{
    "1208/1208.4105.txt": {
        "abstract": "In this paper we investigate the power spectrum of the unresolved  0.5-2 keV CXB with deep {\\em Chandra} 4 Ms  observations in the CDFS. We measured a signal which, on scales $>$30$\\arcsec$, is significantly higher than the Shot-Noise and  is increasing with the  angular scale. We interpreted this signal as the joint contribution of clustered undetected sources like AGN, Galaxies and Inter-Galactic-Medium (IGM).   The power of unresolved cosmic sources fluctuations accounts for  $\\sim$12\\% of  the 0.5-2 keV extragalactic CXB. Overall,  our modeling predicts that $\\sim$20\\% of the unresolved CXB flux is made by low luminosity AGN, $\\sim$25\\% by galaxies and $\\sim$55\\% by the IGM (Inter Galactic Medium). We do not find any direct evidence of the so called Warm Hot Intergalactic Medium (i.e. matter with  10$^5$K$<$T$<$10$^7$K and density contrast $\\delta<$1000), but we estimated that it could produce about 1/7  of the unresolved CXB.  We placed an upper limit to the space density of postulated X-ray-emitting early black hole at z$>$7.5  and compared it  with SMBH evolution models. ",
        "introduction": "The Cosmic X-ray Background (CXB) is the result of a multitude of energetic phenomena occurring in the Universe since the epoch of the formation of the  first galaxies.  Its nature has been investigated in the last 50 years with several telescopes but only in the 80's  it became clear that its main contributors are AGN \\citep{suaeccellenza}. Later it has been found that also  galaxies, galaxy clusters, large scale structures and  diffuse hot gas in the Milky Way are sources contributing to the CXB \\citep{fb}. \\\\ With the launch of ROSAT, {\\em Chandra} and XMM-{\\em Newton},  the knowledge about the nature of  sources contributing to the flux of the CXB became suddenly clear.  Deep surveys like the Chandra Deep Field South \\citep[CDFS,][]{gia,luo,xue} and North \\citep[CDFN,][]{bra}, and the Lockman Hole \\citep{brun}  have almost conclusively resolved the problem of the  CXB below 10 keV. In fact  at the flux limits of  {\\em Chandra} and XMM-{\\em Newton}, about 90-95$\\%$ \\citep{bauer,moretti,lehmer12} of the 0.5-2 keV CXB flux has been resolved\\footnote{In these papers, the fraction of unresolved CXB has been estimated at the flux of the faintest detected source, here we measure the average value on  the investigated area} into point and extended sources. \\\\ CXB synthesis models \\citep[see e.g.][]{tre06,gil07} predict that at   fluxes fainter that the current limits, the AGN  source counts progressively flatten and then galaxies become the more abundant sources.\\\\ Galaxies emerge as the dominant population at faint 0.5-2 keV fluxes, with break-even  point at around $\\sim$10$^{-17}$ \\cgs \\citep[see also][]{lehmer,xue,lehmer12} and can make up for a sizable fraction of the X-ray background (9-16\\%; Ranalli et al. 2005). The contribution from galaxy clusters, down to a mass limit of $\\sim$10$^{13}$ M$_{\\odot}$, has been estimated to be of the order of $\\sim$10\\% of the total 0.5-2 keV  CXB \\citep{gil99,lemze}. \\\\ Several authors studied the clustering properties of the CXB to unveil the nature of the sources producing its flux and their properties \\citep{mart,wu,sol94,bf}. The most used technique is that of the two-point autocorrelation function of the surface brightness of the CXB. While at the time of HEAO-1 the task was to determine what was the contribution of QSO to the CXB, after ROSAT most of the  investigations in this field have been performed to unveil signatures of the Warm Hot Intergalactic Medium \\citep{sol02,sol06,gale} and to test Cosmological models \\citep{diego}. This kind of studies is  therefore extremely powerful to study population of sources which are beyond the resolving and detection capabilities of instruments.   In the GOODS fields, \\citet{HM06, HM07} have shown that, after excising detected point and extended sources plus faint HST detected galaxies, the spectrum of the soft CXB was still showing a signal in the 0.5-2 keV band. While a large amount of it could be attributed to a local component (Solar Wind Charge Exchange and Milky Way thermal emission), they have shown that below 1 keV  the fraction of the resolved CXB is not sensitive to the removal of HST galaxies,  thus arguing for a purely diffuse nature of the unresolved CXB.  In the 1-2 keV band they have shown that, by  removing HST sources areas from the analysis, the fraction of unresolved CXB drops by $\\sim$50\\%. This means that a large fraction of the unresolved 1-2 keV CXB flux could be produced in faint undetected galaxies (active or non-active). Moreover, they speculated that a large fraction of the unresolved CXB may be due to faint 3.6$\\mu$m IRAC sources, suggesting that  high-z or absorbed sources could produce a large part of such a radiation. According to their estimate, the remaining fraction of the CXB (i.e. 2-3\\% of the 0.5-2 keV CXB) remains consistent with the prediction of Warm Hot Intracluster Medium (WHIM) intensity from hydrodynamical simulations \\citep[see e.g][]{cen,ursino,mauro06,ron}.\\\\ In the local Universe about 30-40\\% of  the baryons are missing with respect to what is measured at z$\\sim$3 \\citep{fuku}. Simulations predict that most of these baryons got shock-heated and at z$\\lesssim$1 they should have a temperature of the order 10$^{5}$-10$^{7}$ K and therefore emit thermal X-rays \\citep{cen}.   Controversial evidences on the properties of such a medium have been published so far \\citep[see e.g.,][]{nicastro,kaastra,fang,shull}. Though with  very low surface brightness, the WHIM can be distinguished from other kinds of diffuse emissions on the basis of  its clustering properties \\citep{ursino}.  In fact, its emission follows that of clusters and filaments and peaks at  low redshift (z$\\sim$0.5), thus showing a typical feature in the angular clustering  of unresolved  CXB fluctuations. However, WHIM is not the only expected component of the unresolved CXB arising from thermal emission of the Inter Galactic Medium (IGM). X-ray surveys in the local Universe revealed X-ray emission from local galaxy groups down to masses of the order of 10$^{12}$ M$_\\odot$ \\citep[see e.g.][]{eck}. Since the intensity of the X-ray emission of galaxy groups scales with their mass, a large amount of them has not been detected at moderately high-z. Therefore we also expect a contribution to the overall signal from medium to low mass-groups  at z$>$0.2-0.3.\\\\ Concerning the unresolved extragalactic CXB, it is difficult to distinguish its components with a simple spectral analysis mostly because of the poor energy resolution of X-ray sensitive CCDs (i.e. $\\sim$130 eV, for ACIS-I at 1.5 keV). However, cosmological sources leave a unique imprint in the power spectrum (PS) of the anisotropies of the fluctuations of the unresolved CXB in a way that is related to their clustering and volume emissivity properties. \\\\ The unresolved CXB contains information about all those sources that have not been detected at the deepest  fluxes reachable by deep surveys. Moreover, the amplitude of the PS is not only sensitive to the luminosity density of those sources, but it also provides information about their  bias.\\\\  In this paper we study the power spectrum of the unresolved 0.5-2 keV CXB with the  4Ms CDFS data, the deepest X-ray observation ever performed to date.  We model the anisotropies of unresolved CXB with the state-of-the art results on galaxy and AGN evolution models and observations as well as with modern hydrodynamical cosmological  simulations. \\\\ Throughout this paper we will adopt a concordance $\\Lambda$-CDM cosmology with $\\Omega_m$=0.3, $\\Omega_{\\Lambda}$=0.7, H$_0$=70 h km/s/Mpc and $\\sigma_8$=0.83. Unless otherwise stated, errors are quoted at 1$\\sigma$ level and fluxes refer to the 0.5--2 keV band. We used as reference cumulative 0.5-2 keV CXB flux the recent estimates of \\citet{lehmer12}, S$_{CXB}$(0.5-2)=8.15$\\pm{0.58}\\times$10$^{-12}$ \\cgs~deg$^{-2}$. ",
        "conclusions": "In this paper we presented the measurement of the angular PS of the fluctuations of the unresolved CXB in  the 4Ms observation of the CDFS in the angular range $\\lesssim10\\arcmin$. Poisson noise and spurious signals have been modeled and removed from the measured PS. We  performed a spectral decomposition analysis and showed that after removing the low frequency signal, which can be attributed to the shot noise of unresolved sources that  randomly enter the beam, the amplitude of the fluctuations with extragalactic origin account for $\\sim$12.3\\% of the CXB and the significance of the detection of these cosmic fluctuations is  $>$10$\\sigma$. In the next section we briefly discuss to properties of the populations producing the unresolved CXB fluctuations. \\subsection{The population of  undetected AGN} For AGN we folded the observed evolution scenarios with the population synthesis model of \\citet{gil07}  and  a simple recipe for bias where AGN are tracing DMH with mass Log(M)=13.1 M$_{\\odot}$.  This  population of AGN has a  space density that exponentially declines above z=2.7. The CXB production rate necessary to produce the modeled PS yields to a fraction of the unresolved 0.5-2 keV CXB flux of $\\sim$19\\%.   We predicted a CXB flux produced by undetected AGN of $\\sim$2.0$\\times$10$^{-13}$ \\cgs~deg$^{-2}$.  We also tried to probe different evolution scenarios but our data do not allow to significantly constrain the behavior of the AGN XLF at high-z. In Fig. \\ref{lognlogs} we show the predicted LogN-LogS that, according to our model, satisfies the observed fluctuations compared with recent observations.  \\subsection{The population of undetected X-ray galaxies} Galaxies are  the most numerous population of objects  contributing to the unresolved CXB (see Fig. \\ref{lognlogs}); the power produced by such a population is lower that  that  of AGN since  they are less biased even if they produce more CXB flux. In the soft X-rays, galaxies have been observed up to z$\\sim$1 \\citep[see e.g.][]{lehmer}, and therefore their high-z space density is unknown. We have developed a toy model for the galaxies XLF where their evolution follows that of the star formation in the universe. With this model we estimated that galaxies contribute to $\\sim$25\\% of the unresolved  CXB flux.  In a recent paper, \\citet{dij} used a modeling similar to ours  and found that, in principle, X-ray galaxies could produce all of the unresolved 1-2  keV CXB. However  they did not consider the contribution of other undetected sources that we have shown to produce a large fraction of the unresolved CXB. Overall, the predicted source counts of AGN and galaxies are in good agreement with the measurements of \\citet{xue} and \\citet{lehmer12} in the same field. \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth]{lognlogs.ps} \\caption{\\label{lognlogs}   The 0.5-2 keV logN-logS used to reproduce the observed unresolved CXB fluctuations. In $green$ AGN, $blue$ galaxies and, in  $black-shaded$, the sum of the AGN and galaxies, the uncertainty is attributed to the count-rate to flux conversion.   The $red~shaded~area$ is the logN-logS measured by \\citet{lehmer12} in the CDFS. The $red~line$ represents the upper limit of the logN-logS of miniquasars. The $cyan~shaded~area$ represents the expected counts from the fluctuation analysis of \\citet{mi2}. In the inset we show a zoom onto the miniquasar upper-limit logN-logS. The $blue~triangle$ is the expected number of X-ray sources, at the flux limit of the 4Ms CDFS, produced at z$\\lesssim$10 from collapse of POPIII stars accreting at the Eddington Limit. The $red~triangle$ is number of X-ray sources in case of of direct collapse (Quasi-stars) accreting at $\\lambda_{Edd}$=0.3} \\end{figure}  \\subsection{CXB from IGM and WHIM} On scales larger than 100$\\arcsec$ our analysis shows that the main contribution to  the unresolved CXB fluctuations is due  to  IGM emission  produced by undetected groups, clusters  and the  WHIM. This has been estimated with cosmological hydrodynamical simulations that, together with structure formation,  include feedback mechanisms that pollute the   IGM with metals which are responsible for the X-ray emission  \\citep{ron}. Such a component produces $\\sim$50\\% of the unresolved CXB.\\\\ We determined that $\\sim$1/3 of the IGM flux is produced by the WHIM, where the so called \"missing baryons\"  are expected to lie. However the WHIM definition taken from simulations already applies a density cut ($\\delta < 1000$) that is meant to mimic roughly the effect of excising detected sources. In reality, dense clumps may be present outside virialized objects so we must consider our determination as a lower limit of the total WHIM contribution (see the discussions in Roncarelli et al. 2006, 2012).  Our model thus predicts that the  flux of the WHIM in the CDFS is of the order 1.7$\\times$10$^{-13}$ \\cgs~deg$^{-2}$ (i.e. 2.3\\% of the total CXB flux) and  produces a signal peaking on a few arcmin scale. Such an estimate is about one order of magnitude lower than  what measured by \\citet{gale}  (i.e. 12\\% of the overall diffuse emission in shallower 0.2-0.4 keV XMM-{\\em Newton} observations) where they did not model the contamination from undetected sources. However such an estimate relies on the output of the simulation and is very sensitive to the metallicity of the WHIM. According to \\citet{ron},  different recipes for the metal enrichment of the WHIM may lead to a variation  up to a factor 3 in overall emissivity of the WHIM. More information on such a component of the Universe will be possible if, for example, contamination of undetected cluster could be excised from X-ray maps by masking also optically detected groups with mass Log(M)$>$12-12.5 M$_{\\odot}$. \\subsection{Very high-z sources} Finally, we speculated on   the  possible existence of a population of high-z miniquasars (or early black holes) born from  the collapse of early massive objects. Although we did not detect their signature, we placed an upper limit to their contribution to the CXB  ($<$3$\\times$10$^{-13}$/(1+z) \\cgs~deg$^{-2}$, z$>$7.5). We estimated that these sources would follow the declining  evolutionary track of AGN with Log(L$_X$)$\\lesssim$43.58 erg/s.\\\\ Our observations are not sensitive to the  faint fluxes expected from these sources, and thus we could only place upper limits. However, since such a population peaks at z$>7.5$, their fluctuations should peak on scales of the order of several tens of arcminutes, where the contribution from shot-noise is relatively weak also in shallower surveys. On such a scale,  foreground source population PS significantly  dims and therefore their detection would be possible. To conclude, we determined that at fluxes of the order 10$^{-20}$ \\cgs, their number density is order of $\\sim$60000 deg$^{-2}$ which means that, under the assumption of Euclidean logN-logS,  at the flux limit of the CDFS their number density is of the order 1-2 deg$^{-2}$. \\\\ Using the theoretical  predictions of \\citet{volo}, \\citet{treis} computed the expected number density of the first X-ray sources at z$\\sim$7-10 at the 4Ms CDFS flux limit. In the case of Population III stars remnants accreting at the Eddington limit they find 4.5 deg$^{-2}$, while in the case of direct collapse \\citep[quasi-stars,][]{beg} with accretion regime at $\\lambda_{Edd}$=0.3 their expected number density is  1.9 deg$^{-2}$. In Fig.  \\ref{lognlogs} we show a comparison of these predictions with our upper-limit which, at the zero-th order, is favoring the direct collapse scenario. \\\\  Our  work suggests that future deeper observations on wider fields would allow us to improve the sensitivity of the PS measurement by reducing the Shot-Noise and masking fainter clusters. Moreover a deep survey with a flux limit comparable or deeper to that of the 4Ms CDFS covering  1-2 deg$^{2}$,  would allow us a direct detection of the WHIM feature and, by measuring the large scale shape of the PS, investigate the very  high-z X-ray Universe. A possible future X-ray mission like the proposed Wide Field X-ray Telescope \\citep[WFXT][]{wfxt} will be able to study the PS of the unresolved CXB with high precision given the large collecting area (low photon noise) and the large field of view (low cosmic variance)."
    },
    "1208/1208.3904_arXiv.txt": {
        "abstract": "We examine the projected correlation of galaxies with mass from small scales ($<$few hundred kpc) where individual dark matter halos dominate, out to 15 Mpc where correlated large-scale structure dominates. We investigate these profiles as a function of galaxy luminosity and redshift.  Selecting 0.8 million galaxies in the Deep Lens Survey, we use photometric redshifts and stacked weak gravitational lensing shear tomography out to radial scales of 1 degree from the centers of foreground galaxies. We detect correlated mass density from multiple halos and large-scale structure at radii larger than the virial radius, and find the first observational evidence for growth in the galaxy-mass correlation on 10 Mpc scales with decreasing redshift and fixed range of luminosity. For a fixed range of redshift, we find a scaling of projected halo mass with rest-frame luminosity similar to previous studies at lower redshift. We control systematic errors in shape measurement and photometric redshift, enforce volume completeness through absolute magnitude cuts, and explore residual sample selection effects via simulations. ",
        "introduction": "The presence of dark matter in the Universe is well-established and contributes significantly to structures ranging from galaxies to superclusters. Less understood is the distribution and evolution of dark matter correlated with galaxies over larger volumes beyond the galaxy virial radius.  The WMAP cosmic microwave background (CMB) result for $\\Omega_{\\rm m} = 0.27\\pm 0.03$ \\citep{2011ApJS..192...18K} is over twice that derived from N-body model fits to lensing studies of the inner mass profile of luminous galaxies together with cosmic luminosity density \\citep{2012ApJ...746...38M}.  This is not surprising, since about half the dark matter (DM) is expected to be in virialized halos \\citep{2006ApJ...639..590F}.  Detailed measurements of the mass distribution on large scales around galaxies, or ``galaxy-mass correlation,'' as a function of galaxy properties can thus be a diagnostic of structure formation and evolution.  A universal mechanism for hierarchical structure formation was developed by \\citet{1974ApJ...187..425P}, drawing on bottom-up structure formation ideas of \\citet{1965ApJ...142.1317P} and \\citet{1968MNRAS.141....1S}. More generally, growth of very large scale mass structures is cosmology dependent, and is one of the probes of the physics of dark energy \\citep{2006astro.ph..9591A}.   Weak gravitational lensing (WL) is the only direct probe that can both measure the mass profile associated with galaxies over a wide range of radii crossing the virialization and turnaround scales and does not require assumptions about the dynamical state or baryon content of the system in question.  WL is an inherently statistical technique: many galaxies are required. The dimensionless shear signal ranges from less than 0.1\\% (large-scale structure) to 1\\% (galaxies) to 10\\% (clusters of galaxies).  The WL signal from foreground DM halos lensing more distant galaxies is commonly known as ``galaxy-galaxy lensing'' (GGL) and is measured by cross-correlating the positions of foreground DM halos (as traced by their resident galaxies) with the lensing-induced shear of background galaxies. Most GGL results in the literature focus on the inner parts of the DM halo mass profile inside the virial radius, which is dominated by contributions from individual halos.  Observations of luminosity scaling of the halo mass have been reported by \\citet{2005ApJ...635...73H}, \\citet{2006MNRAS.368..715M}, \\citet{2006A&A...455..441K}, and \\citet{2011arXiv1107.4093V}.  Recent studies have also tackled the problem of redshift evolution of halo mass \\citep{2010ApJ...716.1579L,2012ApJ...744..159L}.  However, the study of GGL outside the virial radius is difficult from both observational and theoretical standpoints.  In observations, residual systematics can masquerade as the lensing signal on large scales \\citep{2005MNRAS.361.1287M}.  On the theory side, we need detailed N-body simulations and ray-tracing studies or comprehensive statistical methods to interpret the signal \\citep[e.g.][]{2005MNRAS.362.1451M,2008MNRAS.388....2H,2012ApJ...746...38M}.   With five 4 square degree deep fields imaged in BVRz$^{\\prime}$, the Deep Lens Survey offers the possibility of measuring galaxy-mass correlations tomographically over a wide range of projected radii well beyond the halo virial radius, and over a useful range of redshift.  In this paper, we investigate how the average galaxy-mass correlation over a wide range of radii varies with luminosity and redshift up to $z$=0.75.  We show the first observational evidence for growth of projected galaxy-mass correlations on 10 Mpc scales.  Section 2 briefly describes the formalism for weak lensing; Section 3 describes the data with particular attention to the measurement of both image shapes and photometric redshifts; Section 4 describes our investigation of photometric redshift error propagation; we present the results in Section 5 and compare them with previous observations and simulations; and we discuss and summarize our results in Section 6.  The Appendix contains further details of systematics tests and correlation matrices.  Throughout, we assume a $\\Lambda CDM$ universe with $H=70$ km Mpc$^{-1} $s$^{-1}$, $\\Omega_{M}=0.27$, and $\\Omega_{\\Lambda}=0.73$.  Distances are given in units of comoving Mpc. ",
        "conclusions": "We have presented projected galaxy-mass correlations measured using galaxy-galaxy lensing over a wide baseline in radius, luminosity, and redshift.  On small radial scales, these signals probe the individual DM halos in which galaxies reside.  On larger radial scales, these signals are sensitive to correlated large-scale structure and neighboring DM halos.  We have focused on two lines of investigation:  fixing the redshift range and varying the mean rest-frame luminosity and conversely, fixing the luminosity range to a volume-complete sample and varying the mean redshift.  For fixed redshift and varied luminosity, we find the well-established scaling of the halo mass with luminosity: intrinsically brighter galaxies are also more massive.  In Figure~\\ref{F:ML_PL}, we show a comparison of our results for the M$_{halo}$-L relation at z$\\sim$0.6 with those of \\citet{2006MNRAS.368..715M} and \\citet{2011arXiv1107.4093V} at z$\\sim$0.1.  After corrections for differing mass definitions and passive luminosity evolution, we find reasonable agreement with the literature results.  On larger radial scales, the shape of the galaxy-mass correlations and their dependence on galaxy luminosity look quite similar to the simulation results of \\citet{2008MNRAS.388....2H}. There are currently few lensing observations of galaxies over wide ranges of luminosity and redshift at radii greater than a few Mpc. We can qualitatively compare our results shown in Figure~\\ref{F:lum} with four observations at lower redshift.  Using shear measurements around z$\\sim$0.1 SDSS galaxies, \\citet{2004AJ....127.2544S} found trends with luminosity that are similar to the higher redshift findings presented here; however, they fit power laws to their shear and galaxy-mass correlations and did not see significant deviation from power laws at large scales.  We can also compare qualitatively to \\citet{2010Natur.464..256R}, who measure shear around luminous red galaxies (LRG) in SDSS out to large radii and model it with a halo model. Their LRG mass overdensity at z=0.3 is consistent with our low redshift z1 and z2 samples at 10 Mpc.  \\citet{2011arXiv1107.4093V} study galaxies in RCS2 with spectroscopic redshifts from SDSS, limiting the lens redshifts to z$\\sim$0.1.  Their Figure 8 shows halo model fits to the lensing signal out to 10 Mpc, which is consistent with the signal we measure.  In these previous studies, the lens galaxies are closest in mass to our L1 sample, but are generally more massive.  Finally, \\citet{2012arXiv1207.1120M} show results for $\\Delta \\Sigma$ for three redshift bins out to z$\\sim$0.5.  However, a direct comparison is difficult because they probe much more massive and more highly biased galaxies (i.e. LRGs).   Most interestingly, our galaxy-galaxy lensing data allows an investigation of the redshift evolution of the projected galaxy-mass correlation on 10 Mpc scales at fixed lens galaxy luminosity. We find evidence for growth over time of the galaxy-mass correlations on large scales from z$=$0.65 to z$=$0.2 as shown in Figure~\\ref{F:gglz1z2L123}. Interpretation of this result in terms of evolution of bias and growth of LSS mass structure with cosmic time requires additional information. The overall effect as a function of lens sample redshift depends on both LSS growth and galaxy bias as a function of mass and type, and an additional probe such as the lens galaxy clustering signal is necessary to disentangle the growth and bias.  We will combine this two-point auto and cross-correlation information for each of the three lens samples with the galaxy-mass correlations reported here in a joint analysis in future work."
    },
    "1208/1208.1306_arXiv.txt": {
        "abstract": "We present {\\it Spitzer} IRS spectra of four carbon stars located in the Galactic Halo and the thick disc.  The spectra display typical features of carbon stars with SiC dust emission and C$_2$H$_2$ molecular absorption.  Dust radiative transfer models and infrared colors enable us to determine the dust production rates for these stars whilst prior CO measurements yield expansion velocities and total mass-loss rates.  The gas properties (low expansion velocities (around 7 km/s) and strong C$_2$H$_2$ molecular absorption bands) are consistent with the stars being metal-poor.  However the dust content of these stars (strong SiC emission bands) is very similar to what is observed in metal-rich carbon stars. The strong SiC emission  may indicate that the carbon stars derive from a metal-rich population, or that these AGB stars produce silicon.  The origin of the halo carbon stars is not known.  They may be extrinsinc  halo stars belonging to the halo population,  they may have been accreted from a satellite galaxy such as the Sagittarius Dwarf Spheroidal Galaxy, or they may be escapees from the galactic disk. If the stars are intrinsically metal-rich, an origin in the disc would be most likely.  If an $\\alpha$-element enhancement can be confirmed, it would argue for an origin in the halo (which is known to be $\\alpha$-enhanced) or a Galactic satellite. ",
        "introduction": "%  Low- and intermediate-mass stars (LIMS, 0.8--8~M$_{\\odot}$ on the main sequence) end their lives with an intense mass-loss episode on the asymptotic giant branch (AGB).  The mass loss from AGB stars is one of the main sources of enrichment of the interstellar medium (ISM) with newly synthesized elements.  In our Galaxy, mass loss from AGB stars produces well over half of the measured dust \\citep{geh89}.  The AGB also dominates the measurable dust production in the LMC and SMC \\citep{mat09,sri09,boy12}.  The mass-loss mechanism is not fully understood, but a two-step mechanism is widely accepted as plausible (e.g., Winters et al. 2000). Pulsation from the star extends the atmosphere, where the material is dense and cool enough to form dust.  Radiation pressure drives the mass loss, accelerating the grains while the gas is carried along by friction.  Dust formation is the key step in the mass-loss process.  Its composition depends on the C/O ratio. If C/O\\,$>$1, i.e. the star is carbon-rich, then most of the oxygen is trapped in volatile CO molecules, which have high dissociation energy. The dust  forms from the remaining refractory elements, and consists  of amorphous carbon and SiC.  For C/O\\,$<$1, the dust will be oxygen-rich, consisting of silicates and aluminium oxides..  Understanding the effect of metallicity on the mass loss is of prime importance to determining the yields of LIMS in different galaxies and understanding the formation of dust in the early Universe. We thus carried out several surveys of mass-losing AGB stars in the Local Group with the {\\it Spitzer} Infrared Spectrograph \\citep[IRS][]{hou04}, targeting the Large Magellanic Cloud \\citep[LMC][]{zij06, mat06,slo08}, Small Magellanic Cloud \\citep[SMC][]{slo06, lag07,slo08}, and several dwarf spheroidal galaxies \\citep{mat07,lag09,slo09,slo12}.  The results show that dust-production rates from oxygen-rich AGB stars are lower in more metal-poor environments, but for carbon stars, metallicity shows no strong influence \\citep{gro07,slo08}. Subtantial dust-production rates are observed for carbon stars with metallicities as low as $\\sim$0.1 Z$_{\\odot}$ \\citep{lag08, slo09,slo12}.   Most of the material ejected via the mass-loss process is in the gaseous form.  The gas mass in the envelope of an AGB star is typically two orders of magnitude higher than the dust mass to which our infrared surveys are sensitive.  To fully characterize the mass loss from AGB stars and its dependence on metallicity, one needs to study the {\\it gas} mass loss in environments with different metallicities.  Carbon stars have recently been discovered in the Galactic Halo. These  are thought to be metal-poor \\citep{ti98, mau04,mau05,mau07}. \\cite{gro97} observed the CO emission from a carbon star in the Halo.  We recently observed a sample of carbon Halo stars in the CO J $= 3 \\rightarrow 2$ transition and found that their expansion velocities were low compared to AGB stars in the Galactic disc with similar infrared colors \\citep{lag10}.  Here, we report new {\\it Spitzer} IRS observations of four of these carbon stars.  We place these observations in the context of our CO observations and additional IRS observations of three other likely carbon Halo stars. ",
        "conclusions": "We presented {\\it Spitzer} IRS spectra of four Galactic disc and halo carbon stars. Typical features of carbon stars with SiC dust emission and C$_2$H$_2$ molecular absorption are seen  in the four spectra. Dust radiative transfer models enabled us to determine the dust-production rates for these stars.  We measured the gas mass-loss rates and expansion velocities from the stars via modeling of the CO J $= 3 \\rightarrow 2$ transition.  The gas-to-dust mass ratio $\\psi$ was then measured by simply dividing the gas and dust mass-loss rates. The measured mass-loss rates and $\\psi$ are in the range 1--4\\,10$^{-6}$M$_{\\odot}$yr$^{-1}$ and 200--1300 respectively.  The stars in the halo have a low expansion velocity (around 7 km/s), strong C$_2$H$_2$ molecular absorption bands,  strong SiC emission and a high gas-to-dust mass ratios.  The strong SiC emission we observe could be due to an overabundance of $\\alpha$-elements (like silicon) in the Halo, that these stars produce silicon or to the fact that those stars metallicities similar to the Galactic disc.  In the latter case, they may have escaped from the disc.  The low expansion velocities and strong C$_2$H$_2$ absorption bands are arguments in favor of the former.  However, explaining the $\\alpha$ enhancement in the Halo is a challenge.  We conclude that these spectra present us with some intriguing contradictions, with the dust properties being similar to metal-rich carbon stars and the gas properties similar to metal-poor carbon stars.   {Those AGB stars could have formed outside of the Halo, as the mass of these carbon stars is inconsistent with a formation in the Halo. The halo carbon  AGB stars we observed could thus have two possible origins. They could have  formed in a galaxy similar to Sgr dSph, orbiting the Milky Way and being gradually stripped by tidal forces or they could have escaped from the Galactic disc.   This work has shown the necessity to combine mid-infrared spectroscopy with millimeter CO observations to fully characterize the dust production {\\it and} gas mass-loss from AGB stars.  Observations of metal-poor AGB stars in Local Group dwarf galaxies are necessary to study the dependence of this mass-loss on metallicity. The Atacama Large Millimeter/submillimeter Array (ALMA) will allow such observations.  Finally near-infrared spectroscopy of the individual stars will be necessary to precisely determine the metallicity and $\\alpha$-element abundances  of the observed stars."
    },
    "1208/1208.0300_arXiv.txt": {
        "abstract": "We present high-resolution optical spectra of the young brown-dwarf eclipsing binary 2M0535$-$05, obtained during eclipse of the higher-mass (primary) brown dwarf.  Combined with our previous spectrum of the primary alone (Paper~I), the new observations yield the spectrum of the secondary alone.  We investigate, through a differential analysis of the two binary components, whether cool surface spots are responsible for suppressing the temperature of the primary.  In Paper~I, we found a significant discrepancy between the empirical surface gravity of the primary and that inferred via fine analysis of its spectrum. Here we find precisely the {\\it same} discrepancy in surface gravity, both qualitatively and quantitatively. While this may again be ascribed to either cool spots or model opacity errors, it implies that cool spots {\\it cannot} be responsible for {\\it preferentially} lowering the temperature of the primary: if they were, spot effects on the primary spectrum should be preferentially larger, and they are not. The \\teff\\ we infer for the primary and secondary, from the TiO-$\\epsilon$ bands alone, show the same reversal, in the same ratio, as is empirically observed, bolstering the validity of our analysis. In turn, this implies that if suppression of convection by magnetic fields on the primary is the fundamental cause of the \\teff\\ reversal, then it cannot be a local suppression yielding spots mainly on the primary (though both components may be equally spotted), but a {\\it global} suppression in the interior of the primary.  We briefly discuss current theories of how this might work. ",
        "introduction": "2MASS 05352184$-$0546085 (hereafter 2M0535) is a very young system located in the Orion Nebula Cluster, and identified by \\citet[][hereafter SMV06]{stassun06} as the first known substellar eclipsing binary (EB).  EBs allow extremely precise direct measurements (via their orbital dynamics and eclipse lightcurves) of the component masses and radii, and hence their surface gravities ($\\log g$), as well as the ratio of their luminosities (and thus the ratio of their effective temperatures, \\teff).  2M0535 therefore enables one to rigorously test the evolutionary models and synthetic spectra that are widely used to characterize the vast majority of brown dwarfs (for which a direct determination of mass and radius is not possible), and to do so at the young ages where our theoretical knowledge is most lacking.  The measurements by SMV06 were refined by \\citet{stassun07} and \\citet[][hereafter G09]{gomez09}. The central results are: {\\it (1)} Both components of 2M0535 are moderate mass brown dwarfs ($M_{\\rm A} = 0.0572 \\pm 0.0033 \\,\\msun$, $M_{\\rm B} = 0.0366 \\pm 0.0022 \\,\\msun$); {\\it (2)} their radii ($R_{\\rm A} = 0.690 \\pm 0.011 \\rsun$, $R_{\\rm B} = 0.540 \\pm 0.009 \\rsun$) are consistent with the theoretical expectation that young brown dwarfs should be much larger than their field counterparts; and {\\it (3)} the \\teff\\ ratio of the components ($T_{\\rm eff,B}/T_{\\rm eff,A} = 1.050 \\pm 0.004$) shows a surprising reversal, with the more massive primary (A) being cooler than the secondary (B).  The \\teff\\ reversal is not predicted by any current set of theoretical evolutionary tracks.  To explain it, \\citet[][henceforth CGB07]{chabrier07} proposed that strong magnetic fields on the primary suppress convection, both globally in the interior, and/or locally near the surface (producing cool surface spots); neither effect is included in standard evolutionary models, and both would act to depress the effective temperature of the primary.  \\citet[][hereafter MM09]{macdonald09} subsequently put forward a qualitatively similar hypothesis, wherein magnetic fields lower the \\teff\\ of the primary by inhibiting interior convection (though their theory differs in important respects from that of CGB07, as we discuss later).  Bolstering the case for strong magnetic fields preferentially affecting the primary, \\citet{reiners07} found that, compared to the secondary, the primary is a relatively rapid rotator (which should boost field generation), with very prominent chromospheric H$\\alpha$ emission (ultimately powered by the release of magnetic stresses). \\citet{mohanty09} subsequently showed, through an analysis of the optical to mid-infrared spectral energy distribution of 2M0535, that ongoing disk accretion is highly improbable in this system. Thus the H$\\alpha$ emission in the primary is indeed likely to be chromospheric in origin, supporting Reiners et al.'s conclusion that it harbors strong magnetic fields.  Interestingly, these results are also consistent with the behaviour of low-mass {\\it field} stars: chromospherically active late-K and M field dwarfs appear cooler than inactive ones of the same luminosity \\citep{morales08,stassun12}.  The implication is that magnetic fields appear to have a large impact on the effective temperature of low-mass stars and brown dwarfs.  If true, this would have far-reaching consequences for our understanding of these objects.  For example, the initial mass function (IMF) in young star-forming regions is often derived by comparing (sub)stellar luminosities and temperatures to theoretical HR diagrams. The latter do not include any field effects, while a substantial fraction of low-mass stars and brown dwarfs at these very early ages show evidence of strong fields (rapid rotation and high activity), just like 2M0535A and active field dwarfs; thus, the inferred IMF might be skewed by a field-induced depression in \\teff, potentially causing an overestimate of the number of low-mass brown dwarfs.  The issue now is to decipher how exactly fields achieve this effect, if indeed they are responsible.  As noted above, theory suggests they might do so by producing cool surface spots (through local suppression of convection), and/or by inhibiting convection globally in the (sub)stellar interior.  Observationally, checking for the presence of spots is clearly easier.  G09 carried out the first investigation of spots on 2M0535AB.  They showed that the small amplitude residual (non-eclipse) variations in the system's lightcurve, modulated at the rotational periods of the primary and secondary, can be well-reproduced by cool spots asymmetrically covering a small fraction ($\\lesssim$ 10\\%) of both components' surfaces.  While they find no evidence for a large ($\\gtrsim$ 50\\%) spot coverage preferentially on the primary, as required by CGB07's theory if spots are to account for the \\teff\\ reversal, they cannot rule out such spots either, if the latter are arranged {\\it symmetrically} about the primary's rotation axis (e.g., polar spots, latitudinal bands, or ``leopard spots'').  To check for spots independent of their orientation, one must examine the spectra of the binary components.  Cool spots are by definition cooler than the surrounding photosphere, and the effective surface gravity within them is also lower (because the magnetic pressure partly offsets the gas pressure, mimicking the reduction in gas pressure caused by a lower surface gravity); both effects alter the shape of temperature- and gravity- (more accurately, gas-pressure-) sensitive spectral lines relative to an unspotted spectrum.  We embarked upon this study in a previous paper \\citep[][hereafter Paper~I]{mohanty10}, where we analysed the high-resolution optical spectrum of the primary 2M0535A alone (obtained during secondary eclipse, with a negligible 1.6\\% contamination by the secondary). Specifically, we derived \\teff\\ and $\\log g$ by simultaneously fitting the observed TiO-$\\epsilon$ band and red lobe of the KI doublet with state-of-the-art synthetic spectra, and compared our results to the empirically known $\\log g$.  We found that at the \\teff\\ $\\approx$ 2500 K inferred from TiO, the KI lobe implies $\\log g$ $\\approx$ 3.0, or 0.5 dex lower than the empirical value.  Conversely, at the known $\\log g$ of 3.5, the \\teff\\ inferred from \\ion{K}{1} is 2650 K, or 150 K higher than derived from TiO.  Such discrepancies are indeed expected if the photosphere is spotted, due to the temperature and effective gravity properties of spots noted above.  In particular, we showed that the spectrum of 2M0535A is consistent with an unspotted stellar photosphere with \\teff\\ = 2700 K and (empirical) $\\log g$ = 3.5, coupled with axisymmetric cool spots that are 15\\% cooler (2300 K), have an effective $\\log g$ = 3.0 (0.5 dex lower than photospheric) and cover 70\\% of the surface. The spot temperature and gravity are consistent with the properties of sunspots and starspots in general, as well as with the previous lightcurve analysis of 2M0535, while the covering fraction agrees with CGB07's requirement if spots are to cause the \\teff\\ reversal.  On the other hand, these discrepancies may arise from errors in the molecular opacities or equation of state (EOS) in the synthetic spectra we used to fit the data.  Such errors are known to be present from analyses of field dwarfs \\citep{reiners05}; while it is unclear whether they persist in the model spectra at the much lower \\logg\\ appropriate for very young brown dwarfs, it is certainly possible (see Paper~I). ",
        "conclusions": "\\subsection{Implications for Spots Causing \\teff\\ Reversal} Assuming that our results are not substantially affected by the lack of high S/N data in the very cores of the lines in the secondary spectrum (as we have argued above they are not), there are only four possible interpretations of the combined analysis presented here and in Paper~I.  \\noindent {\\it (1)}\\,  Spots are the major cause of the spectral discrepancy in 2M0535A and B\\footnote{As an aside, we note that this scenario is unlikely: {\\it Both} components would then have a spot coverage of $\\sim$70\\% (as shown by our analysis of the primary in Paper I), which is an extremely large fraction (effectively making the stellar surface appear to be a very cool one covered with hot spots, instead of the reverse).  While one may entertain such extensive spottedness on one object, it strains belief to consider it on both.}.  In this case, given that we find the same discrepancy, both qualitatively and quantitatively, in both components, spots cannot be responsible for preferentially depressing the \\teff\\ of the primary.  \\noindent {\\it (2)}\\,  Opacity/EOS errors cause most of the discrepancy instead.  If so, then the spot coverage on the primary would be far too small to produce the \\teff\\ reversal: as noted in \\S\\ref{goals}, it is only by ascribing the {\\it entire} discrepancy in the primary to spots, without considering any model errors, that we get the very large coverage required by CGB07's spot theory.  \\noindent {\\it (3)}\\,  Spots and model errors contribute equally to the discrepancy in each component.  In this case, {\\it both} of the above conclusions would hold: the spot coverage on the primary would be too small, and there would anyway be no marked difference in spottedness between the components, again excluding spots as the cause of the \\teff\\ reversal.  \\noindent {\\it (4)}\\,  Spots cause the discrepancy in the primary, but model errors, or some other effect, cause it in the secondary. This is unlikely in the extreme, requiring a monumental coincidence. In particular, the empirical $\\log g$ of 2M0535A and B are nearly identical, and their \\teff\\ are very similar as well (as evinced by the empirical \\teff\\ ratio of $\\sim$1.05, very close to unity). Model errors would then have to be very prominent at the secondary's \\teff\\ and $\\log g$, but disappear over the small parameter jump to the primary, which is improbable; simultaneously, the effects of such errors on the secondary spectrum would have to coincidentally mimic exactly the spot effects on the primary, which is even more unlikely. The same argument applies to any other effect invoked for the secondary but not the primary.  Together, these lines of reasoning imply that, while spots may be present on both 2M0535A and B (and almost certainly are to some extent, as shown by the lightcurve analysis of G09), they {\\it cannot} be responsible for {\\it preferentially} lowering the \\teff\\ of the primary by a large amount, and thus causing the \\teff\\ reversal.  \\subsection{General Implications for Magnetic Fields Causing \\teff\\ Suppression}  In light of this result, one might postulate that magnetic fields are not responsible for the temperature reversal at all, and perhaps heating due to tidal interactions (Heller et al.\\,2010; Gomez Maqueo Chew et al.\\,2012) is to blame instead.  However, with an orbital period of $\\sim$10 days, and rotation periods of $\\sim$3 days and $\\sim$14 days in the primary and secondary respectively (Gomez Maqueo Chew et al. 2009), the two brown dwarfs in 2M0535-05 are sufficiently well separated and sufficiently non-synchronous that significant tidal interactions can be reasonably ruled out (Heller et al.\\,2010, though a more thorough treatment of the coupled evolution of tidal effects and substellar structure is required to confirm this, as the latter authors note).  On the other hand, even if spots, caused by a local suppression of convection by magnetic fields, are ruled out, fields may still produce the \\teff\\ reversal by globally inhibiting convection in the interior of the primary. Both CGB07 and MM09 have proposed such a mechanism.  The two theories differ significantly in the interior field strengths invoked, however.  MM09 apply a modified Gough-Taylor instability criterion, in which the magnetic energy basically scales as fraction ($\\sim$1--10\\%) of the total thermal energy, to derive required field strengths of order 10--100 MG in the interior.  In more recent work, they derive much smaller fields, but still $\\sim$1 MG (MacDonald \\& Mullan 2012).  However, these fields are orders of magnitude greater than the few 10s of kG interior fields, in equipartition with the kinetic (turbulent and convective) energy, suggested by recent simulations \\citep{browning08,browning10}\\footnote{Mullan \\& MacDonald (2012) suggest that the simulations by \\citet{browning08} probe relatively small rotational angular velocities, compared to fast rotating M dwarfs, and thus the field strengths derived in the simulations may be linearly scaled up with rotation rate.  However, the stars in these simulations are in fact quite close to saturation, so it is not clear that a linear increase in field strength with rotation is applicable (M.\\,Browning, in prep.).}.  Moreover, interior fields $\\gtrsim$\\,1\\,MG may also be buoyantly unstable; more detailed simulations are needed to check if this can be avoided.  Conversely, CGB07 present qualitative physical arguments suggesting that interior fields of $\\sim$10\\,kG (consistent with equipartition with the kinetic energy, and with the recent simulational results cited above) can seriously impede global convection in young brown dwarfs.  Within the context of mixing length theory, they find that a convective length scale parameter of $\\alpha \\ll 0.5$ is required to explain the \\teff\\ reversal in 2M0535AB.  Testing this, however, requires detailed full 3D simulations of cooling flows along flux tubes in a magnetized convecting medium, which is a considerable undertaking.  Nevertheless, such simulations are essential to discriminate between the theories above, and test whether magnetic fields can indeed inhibit interior convection sufficiently to explain the observed \\teff\\ reversal.  The rapid pace of advance in simulational complexity \\citep[e.g.][]{browning08,browning10} gives one hope that this will be become a reality in the not-too-distant future.  Concurrently, more observations are required to test the universality of such temperature suppression in young very low mass stars and brown dwarfs, not just in eclipsing systems but also in isolated objects.  The techniques used by \\citet{morales08} for field low-mass stars offer a way forward here, though uncertainties in age-dependent luminosities for young objects will simultaneously present a problem in applying these methods.  However difficult a task, it must be tackled in order to really understand how magnetic fields affect the fundamental parameters of stars and brown dwarfs."
    },
    "1208/1208.2075_arXiv.txt": {
        "abstract": "We present high resolution $H$-band polarized intensity ($PI$; FWHM = $0.''1$: 14 AU) and $L'$-band imaging data (FWHM $= 0.''11$: 15 AU) of the circumstellar disk around the weak-lined T Tauri star PDS~70 in Centaurus at a radial distance of 28 AU ($0.''2$) up to 210 AU (1.$''$5). In both images, a giant inner gap is clearly resolved for the first time, and the radius of the gap is $\\sim$70 AU. Our data show that the geometric center of the disk shifts by $\\sim$6 AU toward the minor axis. We confirm that the brown dwarf companion candidate to the north of PDS~70 is a background star based on its proper motion. As a result of SED fitting by Monte Carlo radiative transfer modeling, we infer the existence of an optically thick inner disk at a few AU.  Combining our observations and modeling, we classify the disk of PDS~70 as a pre-transitional disk. Furthermore, based on the analysis of $L'$-band imaging data, we put an upper limit mass of companions at $\\sim$30 to $\\sim$50$M_{\\rm J}$ within the gap. Taking account of the presence of the large and sharp gap, we suggest that the gap could be formed by dynamical interactions of sub-stellar companions or multiple unseen giant planets in the gap. ",
        "introduction": "Protoplanetary disks are believed to be the birthplaces of planets \\citep[e.g.,][]{haya85}; hence, understanding the evolution of these disks guides our understanding of the process of planet formation. Disks which have substantial infrared excesses but reduced fluxes at wavelengths $\\lesssim$20 $\\mu$m, i.e., transitional disks \\citep{stro89}, could be related to the early phases of planet formation \\citep[see a recent review of][]{will11} and are therefore particularly important for understanding how, where, and when planets form. For many transitional disks, partial inner holes or partial gaps have been directly resolved by interferometry at (sub)millimeter wavelengths \\citep[e.g.,][]{piet06,andr11} and imaging at near-infrared wavelengths \\citep{fuka06,thal10,hash11}. Numerous mechanisms have been proposed to explain the clearing of gaps in transitional disks, including grain growth \\citep[e.g.,][]{dull05}, photoevaporation \\citep[e.g.,][]{clar01}, and gravitational interactions with orbiting planets \\citep[e.g.,][]{papa07,zhu11}. Two possible methods to distinguish the disk-planet interactions from other aforementioned proposed gap-clearing mechanisms could be the detection of (1) a planetary companion in the inner hole/gap region \\citep[e.g.,][]{krau12} or (2) a ring-like gap between optically thick inner and outer disks \\citep[i.e., pre-transitional disk;][]{espa07} because dynamical formation of wide gaps could be the only surviving mechanism for wide gapped disks \\citep[e.g.,][]{papa07,zhu11}  One good candidate to investigate the inner hole/gap region at tens AU in the disk in pre-transitional disks is the weak-lined T Tauri star PDS~70 \\citep[K5 type; 0.82$M_{\\odot}$; $<$10 Myr;][]{greg02,riau06,metc04}. A scattered light disk with a radius at 14 to 140 AU was detected by $Ks$-band imaging \\citep{riau06}. The possible presence of inner and outer disks with different temperatures were suggested by \\citet{metc04} and \\citet{riau06}, which may imply that PDS~70 is a pre-transitional disk object. In this $Letter$, we present high resolution imaging of PDS~70 with Subaru/HiCIAO and Gemini/NICI. ",
        "conclusions": "\\subsection{$H$-band Polarimetry and $L'$-band imaging} Figure~\\ref{f2} shows $H$-band PI images and the $L'$-band LOCI image of PDS~70 assuming a distance of 140 pc, along with radial surface brightness profiles. We find a clear elliptical ring in the $H$-band, which has not been reported in previous high-resolution imaging \\citep{riau06}. A partial elliptical disk is observed in the $L'$-band, due to the inevitable loss of flux in the process of LOCI; hence, we derive radial profiles of suface brightness based on the $H$-band PI image only and companion mass limits from the $L'$-band LOCI image only.  We consider that the ellipse shape is due to the system's inclination, and show the results of fitting an ellipse to these data in Table~\\ref{table1}. The position angle of the major axis and the inclination of the disk are similar with those of $K_{\\rm s}$-band imaging \\citep[PA $\\approx155^{\\circ}$ and $i\\approx62^{\\circ}$;][]{riau06}. We measured an offset of 44 $\\pm$ 3 mas ($\\sim$6 AU) at PA=87.9$^{\\circ}$ between the geometric center of the disk and the central star. The positional accuracy of the central star is 1.5 mas (0.2 AU). The direction of this offset is roughly consistent with that of the minor axis of 68.6$^{\\circ}$, and the sign of this offset indicates that the southwest side is inclined toward us (i.e. the near side) (see the model image in \\citealt{dong12b}). This geometry is also consistent with the facts that (1) the northeast side of the disk is wider due to the back illumination of the wall \\citep{thal10}, and (2) the southwest side is brighter than the northwest side due to forward scattering \\citep{fuka06}.   Assuming the cavity edges correspond to the peak PI of the disk, a radius of the cavity is measured as $\\sim$70 AU. The outer radius of the disk is measured to be $\\sim$140 AU in figure~\\ref{f2}(d), and corresponds to the location at which our sensitivity is no longer sufficient to detect extended emission.  A single power-law fit was performed to the radial profiles along the minor and major axes (figure~\\ref{f2}d and e). Our results of $\\sim$ -3.6 and $\\sim$ -1.7 are different from $\\sim$ -2.8 of \\citet{riau06}.  We also found a flux deficit of $PI$ in the direction of the minor axis. Since the observed scattering angle at the minor axis deviates from 90$^{\\circ}$, the polarization fraction along the minor axis is lower. The $PI$ at the minor axis is therefore lower than at the major axis, and such hole-like structures are similar to those discussed in \\citet{perr09}.  We checked the proper motion of the companion cadidate to the north of PDS~70 reported by \\citet{riau06}.  PDS~70 has a proper motion of $(\\mu_{\\alpha}{\\rm cos} \\ \\delta, \\mu_{\\delta}) = (-24.7 \\pm 11.4, -13.3 \\pm 11.4)$ mas/yr \\citep{rose10}; the separation between PDS~70 and the companion cadidate should increase if the companion cadidate is a background star. Since the separation in HiCIAO and NICI images are 324.13 $\\pm$ 0.15 AU and 324.44 $\\pm$ 0.10 AU, respectively, the separation in \\citet{riau06} is inferred to be $309 \\pm 12$ AU. Our estimation has a good agreement with the actual observed separation of 301.75 $\\pm$ 0.06 AU in \\citet{riau06}, and therefore, we concluded that the companion cadidate is a background star.   \\subsection{Detectable planetary-mass companions} Since the follow-up $L'$-band observations with Gemini/NICI failed to detect any significant signals of point-like sources in the gap we put constraints of upper limits for companion(s). Figure~\\ref{f2}(f) shows the detectable masses of companions at 5$\\sigma$. The LOCI parameter of the optimization area is 250 $\\times$ FWHM, which is different from that described in sec.~\\ref{loci}. For that, we first applied a median filter with 0.$''$11 width to the image, and then calculated the standard deviation as a noise level in concentric annuli along the major axis of the disk. The mass was calculated by assuming the COND evolutionary model \\citep{bara03}, a distance of 140 pc, and an age of 10 Myr \\citep{metc04}. We took into account the flux loss due to the partial self-subtraction by testing how point sources are affected by LOCI. The detectable mass limit is tens $M_{\\rm J}$, therefore, stellar companions down to masses associated with massive brown dwarfs are excluded within the gap.  \\subsection{Modeling of spectral energy distribution (SED)}\\label{model} Although the SED fitting for PDS~70 has been performed in previous studies \\citep{metc04,riau06}, both the availability of new archival photometric data (table~\\ref{table2}) and our imaging results motivates us to revisit the SED of the system using Monte Carlo radiative transfer (MCRT). Note that though PDS~70 has been observed by Spitzer IRS, the object was mispointed by $\\sim$2.3$''$ with components along both the spatial and dispersion axes, therefore, the IRS spectrum of PDS~70 is not used in this work.  {\\bf Setup for modeling.} The method for our MCRT simulations is described in \\citet{dong12a} and \\citet[in prep.]{whit12}. In a subsequent paper II \\citep[in prep.]{dong12b}, we will perform a detailed radiative transfer modeling of both the SED and the SEEDS imagery, and present a fiducial disk+cavity model which reproduces both observations well. We will also explore the parameter space around the fiducial model in paper II, and provide a full discussion on the constraints on the various model parameters there. In this letter, we only briefly describe the fiducial model and its resulting SED. We note that this model is not a {\\it best} fitting model in an absolute sense, since a full $\\chi^2$ fitting of the observations with all the free parameters in a protoplanetary disk is essentially impossible \\citep{math12}. However the constraints on many parameters such as the cavity size, depletion, and the surface density of the inner disk are reasonably tight, as will be shown in paper II.  Our model contains a cavity 70 AU in radius. The surface density both inside and outside the cavity decreases with radius as $\\Sigma\\propto \\frac{R_c}{R}e^{-R/R_c}$, where $R_c=50$ AU, while $\\Sigma$ inside the cavity is reduced to $\\delta\\times$ the extrapolated value from the outer disk, with $\\delta$ being the depletion factor. The temperature structure of the disk is determined from the radiative transfer calculations. The inner edge of the disk is self-consistently determined at the dust sublimation temperature ($\\sim1600$ K). We ignore accretion in the model, as suggested by its nature of being a weak line T Tauri star \\citep{greg02}. We use a pre-main sequence star of spectral type K5, radius 1.39 $R_\\odot$, mass 0.82 $M_\\odot$, and temperature 4500~K for the central source, as suggested by \\citet{greg02} and \\citet{riau06}. The disk has a gaussian density distribution in the vertical direction, with scale heights $h$ as input parameters. Two disk components are included: a thick disk with small grains (sub-micron-sized), representing the pristine ISM-like dust; and a thin disk with large grains (up to $\\sim$mm-sized) and 20\\% of the scale height of the small grains, as the result of grain growth and settling \\citep{dull04a,dull04b,dull05}. Most of the dust mass (0.967) is in the settled disk, and the total dust-to-gas ratio is assumed to be 1/100. The SED is produced assuming a disk inclination angle $50^\\circ$. The other parameters are summarized in table~\\ref{table1}.  {\\bf SED fitting.} Figure~\\ref{f3} shows the good agreement between our model SED and available photometric data. As we will show in paper II, the thermal emission from the cavity wall at $\\sim$70 AU peaks at $\\sim40\\mu$m. The wall emission is a major signature of (pre-)transitional disk SED \\citep[e.g.,][]{espa07}.  The surface density of small dust is $\\sim0.001$g cm$^{-2}$ at 0.1 AU, and the opacity of small dust is $\\sim10^4$cm$^2$ g$^{-1}$ at $\\sim1\\mu$m (roughly the peak of the stellar radiation). This makes the inner disk optically thick in the vertical direction. Since we assume a surface density profile decreasing with increasing radius, when moving out the cavity gradually becomes vertically optically thin, and the transition happens at $\\sim1$ AU if only taking into account the small dust (the existence of big dust inside the cavity is poorly determined from SED and scattered light image). As we will discuss in detail in paper II, this optically thick inner disk is needed to explain the $\\sim$2-40 $\\mu$m SED, as models with optically thin inner disk fail to reproduce the SED. Following the convention in the literature \\citep{espa07}, we classify PDS~70 as a pre-transitional disk object.  \\subsection{Origin of the gap}  Grain growth is one mechanism which can potentially form disk gaps \\citep[e.g.,][]{birn12}. It is capable of reproducing the SED of transitional disks but not observed millimeter-wavelength images. However, the sharpness of PDS~70's $H$-band gap edge suggests grain growth is unlikely to be the primary reason, since the models generally predict smooth gap edge features \\citep[e.g.,][]{birn12}.  Photoevaporation is another possible mechanism. A photoevaporative wind can prevent outer disk material from feeding the inner disk; without this supply, the inner disk material will rapidly accrete onto the central star, creating a cavity inside out \\citep{clar01}. Although a pre-transitional disk-like structure could be produced for a very short period \\citep[see figure~1 in][]{alex06}, it is quite unlikely to observe such a snap-shot by coincidence, so that it may be difficult to explain the existence of the optically thick inner disk inferred from the SED.  Dynamical interactions with (sub)stellar companions \\citep[e.g.,][]{arty94} or orbiting planets \\citep[e.g.,][]{papa07,zhu11} are also potential mechanisms, as several binary systems \\citep[e.g., CoKu Tau 4;][]{irel08} are known to have a cavity in their circumbinary disk, and simulations show that multiple planets can open wide gaps \\citep{zhu11}, Furthermore, the survival of an optically thick inner disk within a few AU in \\citet{zhu11} mimics the inner disk of PDS~70. Our $L'$-band observation has only ruled out companions with a mass of $\\gtrsim$50$M_{\\rm J}$ in the gap region. Therefore, dynamical interactions with low mass brown dawarfs or giant planets may be the origin of the gap. To put a futher constraint of an upper limit mass, future observations should be pursued, such as the aperture-masking interferometric observations \\citep[e.g.,][]{naka89}, which has a better contrast than LOCI imaging. Since PDS~70 harbors the pre-transitional disk with one of the largest gaps, this object may be one of the best candidate for future high-contrast planet imagers.  \\bigskip We are grateful to an anonymous referee for providing many useful comments leading to an improved paper. We appreciate support from the Subaru Telescope and the Gemini South Telescope staff, especially from Jennie Berghuis and Dr.~Tom Hayward. This work is partly supported by a Grant-in-Aid for Science Research in a Priority Area from MEXT, by the Mitsubishi Foundation, and by the U.S. National Science Foundation under Awards No. 1009203 and 1009314.  \\clearpage  \\begin{table} \\begin{center} \\caption{The results of the ellipse and the SED fitting for the disk of PDS~70\\label{table1}} \\begin{tabular}{cc} \\tableline \\tableline \\multicolumn{2}{c}{Ellipse fitting\\tablenotemark{1}} \\\\ \\tableline Diameter of the major axis\\tablenotemark{2} (AU)               & 137.0 $\\pm$ 1.1 \\\\ Diameter of the minor axis\\tablenotemark{2} (AU)               &  88.6 $\\pm$ 0.5 \\\\ Position angle of the major axis ($^{\\circ}$)  & 158.6 $\\pm$ 0.7  \\\\ Inclination by ellipse fitting\\tablenotemark{3} ($^{\\circ}$)      &  49.7 $\\pm$ 0.3  \\\\ Geometric center\\tablenotemark{2,4} (AU : AU ) & ($6.1$ $\\pm$ 0.3 : 0.2 $\\pm$ 0.4) \\\\ \\tableline \\multicolumn{2}{c}{SED fitting\\tablenotemark{5}} \\\\ \\tableline $M_{\\rm disk}$ & 0.003 \\\\ $f$ & 0.967 \\\\ $h^{\\rm o}_{\\rm b}$(100AU) & 2         \\\\ $b^{\\rm o}_{\\rm b}$ & 1.2       \\\\ $h^{\\rm c}_{\\rm b}$(100AU) & 2         \\\\ $b^{\\rm c}_{\\rm b}$ & 1.2       \\\\ $\\delta_{\\rm cav,b}$ & $10^{-3}$ \\\\ $h^{\\rm o}_{\\rm s}$(100AU) & 10        \\\\ $b^{\\rm o}_{\\rm s}$ & 1.2       \\\\ $h^{\\rm c}_{\\rm s}$(100AU) & 10        \\\\ $b^{\\rm c}_{\\rm s}$ & 1.2       \\\\ $\\delta_{\\rm cav,s}$ & $10^{-3}$ \\\\ \\tableline \\end{tabular} \\tablenotetext{1}{ In the ellipse fitting for the disk, the peak positions were first directly determined by the radial profile at position angles every 5$^{\\circ}$ the disk. We then conducted an ellipse fit using an implementation of the nonlinear least-squares Marquardt-Levenberg algorithm with five free parameters of lengths for the major and minor axes, position angle, and central positions. } \\tablenotetext{2}{ We assume that the distance of PDS~70 is 140 pc. } \\tablenotetext{3}{ Derived from the ratio of the major and minor axes. } \\tablenotetext{4}{ Central position (0, 0) is corresponding to the stellar position. } \\tablenotetext{5}{ Row (1): Total mass of the disk (assuming a dust-to-gas-ratio 100). Row (2): Mass fraction of big dust in total dust. Rows (3), (5), (8), and (10): Scale height $h$ at 100 AU. Subscripts $b$ and $s$ indicate big dust and small dust, respectively. Superscripts $o$ and $c$ indicate outer disk and cavity, respectively. Rows (4), (6), (9), and (11): Power index $b$ in $h \\propto R^{b}$ in various disk components. Row (7) and (12): Depletion factor of the big and small dust disk. } \\end{center} \\end{table}  \\clearpage  \\begin{table} \\begin{center} \\caption{Archival photometry data for PDS~70\\label{table2}} \\begin{tabular}{lccc} \\tableline \\tableline Wavelength                         & $F_{\\nu}$ (mJy)    & Note          & Plotted color\\\\ &                    &               & in fig.~\\ref{f3} \\\\ \\tableline $U$\\tablenotemark{a,b}             &     9.7            & \\citet{greg92} & cyan    \\\\ $B$\\tablenotemark{a,b}             &    41.8            & \\citet{greg92} & cyan    \\\\ $V$\\tablenotemark{a,b}             &    99.2            & \\citet{greg92} & cyan    \\\\ $R$\\tablenotemark{a,b}             &   160.9            & \\citet{greg92} & cyan    \\\\ $I$\\tablenotemark{a,b}             &   216.3            & \\citet{greg92} & cyan    \\\\ 2MASS ($J$)\\tablenotemark{a,b}     &   311.9 $\\pm$  6.9 & \\citet{cutr03} & magenta \\\\ 2MASS ($H$)\\tablenotemark{a,b}     &   342.7 $\\pm$ 12.6 & \\citet{cutr03} & magenta \\\\ 2MASS ($Ks$)\\tablenotemark{a,b}    &   275.6 $\\pm$  5.8 & \\citet{cutr03} & magenta \\\\ WISE (3.4 $\\mu$m)\\tablenotemark{b} &   188.1 $\\pm$  4.0 & \\citet{cutr12} & green   \\\\ WISE (4.6 $\\mu$m)\\tablenotemark{b} &   142.0 $\\pm$  2.6 & \\citet{cutr12} & green   \\\\ WISE (12 $\\mu$m)\\tablenotemark{b}  &   153.9 $\\pm$  2.3 & \\citet{cutr12} & green   \\\\ WISE (22 $\\mu$m)\\tablenotemark{b}  &   341.8 $\\pm$  0.7 & \\citet{cutr12} & green   \\\\ AKARI (9 $\\mu$m)                   &   201.2 $\\pm$ 25.8 & VizieR II/297  & black   \\\\ AKARI (18 $\\mu$m)                  &   209.8 $\\pm$ 13.4 & VizieR II/297  & black   \\\\ AKARI (90 $\\mu$m)                  &   851.1 $\\pm$ 62.6 & VizieR II/298  & black   \\\\ IRAS (12 $\\mu$m)                   &   251   $\\pm$ 25.1 & \\citet{mosh89} & blue    \\\\ IRAS (25 $\\mu$m)                   &   348   $\\pm$ 27.8 & \\citet{mosh89} & blue    \\\\ IRAS (60 $\\mu$m)                   &   884   $\\pm$ 61.9 & \\citet{mosh89} & blue    \\\\ MIPS (24 $\\mu$m)                   & 349.7\\tablenotemark{c}   $\\pm$ 7.0  & Spitzer Heritage Archive\\tablenotemark{e} & red \\\\ MIPS (70 $\\mu$m)                   & 1049.9\\tablenotemark{c}  $\\pm$ 19.9 & Spitzer Heritage Archive\\tablenotemark{e} & red \\\\ MIPS (160 $\\mu$m)                  & 873.0\\tablenotemark{c,d} $\\pm$ 36.2 & Spitzer Heritage Archive\\tablenotemark{e} & red \\\\ \\tableline \\end{tabular} \\tablenotetext{a}{ An extinction law was adopted from \\citet{math90}. } \\tablenotetext{b}{ Absolute flux conversions in optical, 2MASS, WISE photometric data were adopted from \\citet{bess98}, \\citet{cohe03}, and \\citet{jarr11}, respectively. } \\tablenotetext{c}{ Since photometric data were not available, we conducted aperture photometry for archival images. } \\tablenotetext{d}{ Since the sky value was not sufficiently measured due to a small field-of-view of available data, our photometry may be less reliable. } \\tablenotetext{e}{ This work is based in part on observations made with the Spitzer Space Telescope, obtained from the NASA/ IPAC Infrared Science Archive, both of which are operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with the National Aeronautics and Space Administration. } \\end{center} \\end{table}  \\clearpage  \\begin{figure} \\epsscale{1} \\plotone{fig01.eps} \\caption{$H$-band polarization vectors of PDS~70 are superposed on the $PI$ image with a software mask with 0.$''$4 diameter before subtracting polarized halo (a) and after subtraction (b). The plotted vectors are binned with spatial resolution. The field of view (FOV) is 3.$''$0 $\\times$ 3.$''$0. All plotted vectors' lengths are arbitrary for the presentation purpose. \\label{f1}} \\end{figure}  \\clearpage  \\begin{figure} \\epsscale{1} \\plotone{fig02.eps} \\caption{(a): $H$-band $PI$ image of PDS~70 with a software mask with 0.$''$4 diameter. (b): same with (a), but its features. The solid ellipse indicate the ring-like disk. The filled circle represent the geometric center of the disk. (c): $L'$-band LOCI image of PDS~70 with a software mask with 0.$''$4 diameter. The parameters in LOCI reductions are described in sec.~\\ref{loci}. The FOV of three images are 3.$''$0 $\\times$ 3.$''$0 with a convolution of a spatial resolution. (d) and (e): Radial profiles at yellow hatched regions of minor and major axes in (b). The values of the profile at northwest and southwest are multiplied by ten for the presentation purpose. (f): Detectable mass at 5$\\sigma$ based on the $L'$-band LOCI image. The LOCI parameters are same with those described in sec.~\\ref{loci}, but the optimization area is 250~$\\times$~FWHM ($N_{\\rm A} =$ 250). \\label{f2}} \\end{figure}   \\clearpage  \\begin{figure} \\epsscale{1} \\plotone{fig03.eps} \\caption{Pre-transitional disk model of PDS~70. Filled circles represent archival photometry (see table~\\ref{table2} for photometry data). The solid black line is the best-fit model with a gap of $\\sim$70 AU (see table~\\ref{table1} for model parameters). Separate model components are as follows: stellar photosphere (green dotted line), thermal emission (blue dotted line), and scattered light (red dotted line). \\label{f3}} \\end{figure}  \\clearpage"
    },
    "1208/1208.3465_arXiv.txt": {
        "abstract": "We use a combination of reverberation mapping data and single-epoch spectra of the \\civ\\ emission line in a sample of both low and high-redshift active galactic nuclei (AGNs) to investigate sources of the discrepancies between \\civ- and \\Hbeta-based single-epoch black hole mass estimates.  We find that for all reverberation mapped sources, there is a component of the line profile that does not reverberate, and the velocity characteristics of this component vary from object-to-object.  The differing strength and properties of this non-variable component are responsible for much of the scatter in \\civ-based black hole masses compared to \\Hbeta\\ masses.  The \\civ\\ mass bias introduced by this non-variable component is correlated with the shape of the \\civ\\ line, allowing us to make an empirical correction to the black hole mass estimates.  Using this correction and accounting for other sources of scatter such as poor data quality and data inhomogeneity reduces the scatter between the \\civ\\ and \\Hbeta\\ masses in our sample by a factor of $\\sim$2, to only $\\sim$0.2 dex.  We discuss the possibility that this non-variable \\civ\\ component originates in an orientation-dependent outflow from either the proposed broad line region (BLR) disk-wind or the intermediate line region (ILR), a high-velocity inner extension of the narrow line region (NLR). ",
        "introduction": "\\label{S_Intro}  Observations of galaxies in the local universe have led to the nearly universal acceptance that massive galaxies house a supermassive black hole (BH) at the center of their gravitational potential well.  In the local universe, the masses of the BHs are strongly correlated with some properties of their host galaxies, leading to the $M_{\\rm BH}-\\sigma_*$ and $M_{\\rm BH}-L_{\\rm bulge}$ relationships \\citep[e.g.][]{Ferrarese00, Gebhardt00a, Tremaine02, Marconi03, Graham07, Gultekin09, Graham11}. Because of the apparent tight evolutionary relationship between the BH and its host, one way in which to probe galaxy evolution is to trace the growth of the central BH.  In order to do this, we need reliable methods for measuring BH masses in the near {\\it and} distant universe.  Direct dynamical BH mass measurement methods work only for nearby quiescent galaxies in which the gravitational sphere of influence of the BH can be spatially resolved (see, e.g., the review by \\citealp{Ferrarese05}, though also see \\citealp{Gultekin09} for a discussion of deriving masses without resolving the sphere of influence).  On the other hand, reverberation mapping \\citep{Blandford82, Peterson93} can be applied to active galactic nuclei (AGNs) to measure BH masses both locally and at cosmological distances. Unfortunately, the required length of the spectrophotometric monitoring campaigns needed to make the measurements becomes longer with increasing source luminosity and distance.  As a result, current measurements only extend to sources with $z \\lesssim 0.3$ (see, e.g., \\citealp{Peterson04, Bentz09lamp, Denney10, Grier12b}; see Kaspi et al.\\ 2007 for a tentative result for a $z\\sim$2 quasar).  Extending direct mass measurements to large samples of objects at high redshifts in order to study the co-evolution of BHs and galaxies is not likely to be possible in the near future with current resources. However, the results from reverberation mapping of local AGNs can be used to calibrate {\\it indirect} estimates of BH masses for large samples of objects using single spectra.  These ``single-epoch'' (SE) mass estimates rely on the observed tight correlation between the monochromatic AGN luminosity, $L$, and the radius of the broad line region (BLR), $R$, measured in reverberation mapping experiments.  With this method, only two observables are then needed: (1) a measurement of the AGN luminosity to use as a proxy for the BLR radius through this $R-L$ relationship \\citep{Kaspi05, Kaspi07, Bentz06a, Bentz09rl}, and (2) a measurement of a broad emission-line width (typically H$\\alpha$, H$\\beta$, \\mgii\\,$\\lambda1549$, or \\CIV) to estimate the Doppler-broadened BLR line-of-sight (LOS) velocity dispersion, $V$. Combining these observables, a virial BH mass can be calculated as $M_{\\rm BH}=fRV^2/G$, where $G$ is the gravitational constant and $f$ is a scale factor of order unity related to the BLR geometry and kinematics\\footnote{For reverberation mapping results that use the line width measured from the rms spectrum, an ensemble average, $\\langle f \\rangle$, sets the AGN BH mass scale zero-point, and is determined assuming AGNs follow the same $M-\\sigma_*$ relationship as quiescent galaxies \\citep{Onken04, Woo10}.  For SE masses, $f$ has a different calibration based on the differences between the SE and rms line widths for H$\\beta$ \\citep[see][]{Collin06}.}.  The majority of reverberation mapping results to date use time delays measured from \\Hbeta\\ emission line variability.  Consequently, the most robust SE mass scaling relation is based directly on the $R-L$ relationship for this emission-line \\citep[see][]{Collin06, Vestergaard06, Bentz09rl}.  Scaling relationships for other broad emission lines are also available \\citep{McLure02, Vestergaard06, McGill08, Vestergaard&Osmer09}, but these SE mass scaling relationships are calibrated to the RM-based \\Hbeta\\ mass scale rather than by direct measurements of the time delays for these lines.  In this way, SE masses can be estimated for large samples of AGNs observed in spectroscopic surveys and across a large range of redshifts \\citep[see, e.g.,][]{Kollmeier06, YShen08, Shen11}.  This suite of intercalibrated AGN BH mass scaling relations initially seems the ideal method for using the myriad of optical survey spectra of AGNs to study BH-galaxy coevolution at all redshifts.  However, there are significant concerns as to the reliability of masses based on lines other than \\Hbeta\\ due to observed systematics in these masses compared to \\Hbeta\\ masses.  For example, only tentative reverberation results exist for \\mgii\\ \\citep{Metzroth06}, so the use of \\mgii\\ as a virial mass indicator is based only on its similar ionization potential and line width compared to \\Hbeta\\ \\citep{McLure04, Vestergaard&Osmer09}. Studies have shown, however, that \\mgii-based masses suffer from clear systematics, some of which may be related to the Eddington ratio \\citep{Onken08} or systematic line width differences between \\mgii\\ and \\Hbeta\\ \\citep[see][]{Croom11}.  \\mgii-based masses are also susceptible to biases related to the prescription used for measuring the line width \\citep{Rafiee&Hall11}.  The same is true of SE masses based on \\civ.  The current \\civ\\ mass scaling relationship \\citep[][hereafter VP06]{Vestergaard06} is based on a direct calibration to \\Hbeta\\ reverberation-mapped AGNs.  For this sample, VP06 found a relatively small scatter in the SE masses compared with the RM masses, $\\sim 0.3$ dex, which has since been cited as the typical assumption of the accuracy of SE masses.  Additional studies have also found a general consistency between \\civ-based and Balmer-line-based masses \\citep[e.g.,][]{Greene10, Assef11}.  However, other authors have questioned the reliability of \\civ\\ as a result of finding little consistency or correlation and a large scatter between \\civ- and \\Hbeta-based SE masses.  These studies contend that \\civ-based masses have too much scatter compared with \\Hbeta\\ masses to be a reliable virial mass indicator \\citep[e.g.,][]{Baskin&Laor05, Netzer07, Sulentic07, Shen12}.  One common postulate for this seeming unreliability is that the \\civ\\ emission region is susceptible to outflows and winds \\citep[see][and references therein]{Richards11}.  Such non-virial gas velocities are then believed to bias the resulting SE mass measurements, rendering them unreliable.  Interestingly, however, such dynamics do not seem to affect the \\civ\\ RM-based masses.  In the few objects where RM measurements have been made for multiple emission lines, \\civ\\ results follow the expected virial relation with other low-ionization lines \\citep{Peterson00a} and the BH masses derived from the individual lines, including \\civ, are mutually consistent \\citep{Peterson04}.  Here, we explore this apparent contradiction --- why would RM-based \\civ\\ masses behave as expected but SE masses do not?  In \\S \\ref{S_SEmassAssumptions} we first review the fundamental assumptions that make estimating a SE mass possible.  In \\S \\ref{S_RMmnrms} we investigate properties of the \\civ\\ profile in the RM sample, identifying a source of bias related to the \\civ\\ line profile.  In \\S \\ref{S:shapebias} we quantify this bias and empirically fit a correction to \\civ\\ masses that significantly reduces the scatter between \\civ\\ and \\Hbeta\\ SE masses.  Finally, in \\S\\S \\ref{S_discussion} and \\ref{S_conclusion} we put our results in context to other studies, suggest a physical interpretation to explain our observations in terms of an orientation-dependent outflow, and summarize our findings. ",
        "conclusions": "\\label{S_discussion}  \\subsection{Other Line Width and Shape Characterizations}  The FWHM is the most widely-used line-width characterization in the literature, and it has the advantages of being easy to measure and relatively insensitive to blending in the line wings.  On the other hand, measurements of $\\sigma_l$ are robust for a wide range of line profiles and have been found to be a less biased line characterization in some cases \\citep[see, e.g.,][]{Peterson04, Denney09a, Rafiee&Hall11}.  The problems with measuring $\\sigma_l$ include difficulties in accurately defining the line wings, particularly in low $S/N$ spectra, and in the presence of other blended emission-line features.  Despite these concerns, Figures \\ref{Fig:widthmnvsrms} and \\ref{Fig:massresidVshape}, respectively, show that characterizing the \\civ\\ width with $\\sigma_l$ leads to (1) velocities that are more consistent with the 1:1 relation between the mean and rms spectrum in the reverberation-mapped sample, and (2) masses that show less bias with the line shape, $S$. In fact, when we calculate the \\civ\\ mass using $\\sigma_l$ (i.e., Equation 8 of VP06), the scatter in the mass residuals about the 1:1 relation of 0.28 dex is only slightly larger than the scatter in the corrected FWHM-based masses (0.26 dex) shown in Figure \\ref{Fig:origAcorrmassresid}, albeit with an additional systematic offset of $-0.17$ dex indicating a zero-point calibration difference between the full sample and VP06.  This offset could be due to the \\civ\\ shape bias or other systematic uncertainties in the exact prescription for calculating $\\sigma_l$, such as blending with the red shelf \\citep[see, e.g.][]{Fine10, Assef11}.  In either case, if we adjust the zero-point of the VP06 relation to eliminate the offset, the remaining scatter about unity is only 0.26 dex, equivalent to the FWHM-corrected \\civ\\ masses.  Because our shape correction to the \\civ\\ mass is a combination of only the FWHM and $\\sigma_l$, expanding Equation (1) shows that we have effectively fit a virial relation of the form $M \\propto {\\rm FWHM}^x \\sigma_l^{y}$, where $y=2-x$ in general, and here, $y=1.63$ and $x=0.37$.  The preferentially higher weight of the line dispersion in this expansion explains the similarity of the scatter between the shape-corrected masses and $\\sigma_l$-based masses.  We also determined the \\civ\\ shape correction strictly from the kurtosis of the \\civ\\ profile.  Results are not shown because they were only based on the N07 and A11 samples for which we had access to the data, but are consistent with using $S = {\\rm FWHM}/\\sigma_l$.  This demonstrates that utilizing either (1) a combination of line width plus shape, or (2) a more complex (i.e., than FWHM) characterization of the \\civ\\ line profile is necessary to properly calibrate \\civ\\ and \\Hbeta\\ masses due to the contamination of the non-variable component of \\civ. \\citet{Baskin&Laor05} also found that the scatter between \\civ\\ and \\Hbeta\\ masses could be reduced by fitting a \\civ\\ mass correction based on additional \\civ\\ profile information, albeit with more parameters, to what we have done here.  Additionally, \\citet{Wang09} and \\citet{Park11} have suggested relaxing a strict virial requirement ($M \\propto V^2$) altogether as a means to reduce scatter in the calibration of SE mass scales due to systematic differences between various SE line profiles and the rms \\Hbeta\\ profile to which these mass scales are ultimately calibrated.  \\citet{Rafiee&Hall11} found that this relaxation could also explain the ``sub-Eddington boundary'' for FWHM-based BH masses, but then again, so could using $\\sigma_l$-based masses with the typical virial requirement.   \\subsection{Sources of Remaining Scatter} \\label{S:remainingscatter}  While the \\civ\\ shape correction defined by Equation (1) significantly reduces (by 0.13 dex) the scatter in the \\civ-to-\\Hbeta\\ mass residuals, a scatter of $\\sim$0.3 dex still remains.  We briefly investigated three potential sources for this remaining scatter: \\begin{enumerate}  \\item {\\bf Data quality:} Several absorbed N07 objects (gray stars) became even larger outliers after the \\civ\\ shape correction.  This suggests that the {\\it measured} line width in these objects is not the {\\it intrinsic} line width, so the correction to these sources was meaningless.  This is a serious concern for SE masses measured from survey quality spectra, which are typically of relatively low $S/N$ \\citep[see][and Denney et al., in prep.\\ for additional discussions and investigations of this concern]{Denney09a, Assef11}.  \\item {\\bf Data Inhomogeneity:} As there is no universal prescription for AGN spectral decomposition and line width measurements, inhomogeneity in the data analysis and computational methods used to derive SE masses leads to measurable systematics when trying to compare masses between samples and emission lines.  This is demonstrated by the comparison of our FWHM measurements of the N07 sample in Figure \\ref{Fig:usVnetzerFWHM} \\citep[see also][for additional sample to sample comparisons]{Assef11, Vestergaard11}. There are also no universal mass scaling relationships for each emission line, and differences in the mass scale also lead to differences in masses derived by different studies \\citep[see discussions by][]{McGill08, Shen12}.  \\item {\\bf H\\boldmath$\\beta$ Systematics:} Literature comparisons between SE \\civ\\ and \\Hbeta\\ masses typically blame \\civ\\ masses for any observed inconsistencies, since \\Hbeta\\ is arguably the most well-characterized and best-studied line in terms of direct RM-based masses for comparison.  However, \\Hbeta\\ it is not without systematics of its own, particularly contamination by NLR emission and host galaxy starlight, which are relatively unimportant for \\civ.  The systematics introduced into the \\Hbeta\\ line width and/or optical luminosity can be significant \\citep{Denney09a, Bentz09rl, Park11}.  \\end{enumerate}  We tested if the scatter in the \\civ-to-\\Hbeta\\ masses could be further reduced by addressing these additional sources of scatter.  First, we removed (1) single VP06 epochs with $S/N < 10$ pixel$^{-1}$ in the continuum, (2) the seven objects from the N07 sample that show evidence for absorption (we could not cull $S/N < 10$ pixel$^{-1}$ data from this sample without eliminating it completely), and (3) the three A11 sources defined as Group II objects due to unreliable \\Hbeta\\ widths and other issues \\citep[see][for details]{Assef11}.  Next, we removed host starlight from the VP06 sample by replacing the SE \\Hbeta\\ masses with the RM \\Hbeta\\ masses.  Ideally, host starlight should be removed from all samples, but the required host fluxes are not available.  In any case, the VP06 sample is the most susceptible to this bias because of its relatively lower mean AGN luminosity.  This choice removes some of the consistency between the treatment of this sample compared to the A11 and N07 samples, but these differences are smaller than the accuracy gained by this correction.  Finally, fully removing the data analysis inhomogeneity is not possible. However, we fit and defined a shape-based \\civ\\ mass correction independently for each sample, since the sample-to-sample inhomogeneities in the spectral fitting and line width measurement techniques have likely led to a sub-optimal \\civ\\ shape correlation fit based on the combined sample.  We list the individual sample fit parameters and uncertainties in Table \\ref{Tab_Fits}. Figure \\ref{Fig:bestorigAcorrmassresid} shows the results of these modifications to our analysis.  The top panel again shows the ratio of the \\civ-to-\\Hbeta\\ masses of this ``high-quality'' sample before correcting \\civ\\ masses for the shape bias.  Here the scatter around the 1:1 relation is now 0.33 dex, and the systematic offset is small, showing that simply removing poor quality data and host galaxy starlight alone can reduce the scatter.  The bottom panel shows the residuals after correcting the \\civ\\ masses using the individual shape corrections determined for each homogeneous sample.  Here the scatter is only 0.22 dex, demonstrating the \\civ\\ and \\Hbeta\\ masses to be in very good agreement.  The remaining scatter is less than the previously estimated intrinsic systematic uncertainties of $\\sim 0.3$ dex inherent in SE masses (VP06) and is comparable to the uncertainty in the calibration of the $R-L$ relationship (\\citealp{Bentz09rl}, though see also \\citealp{Peterson10} for evidence that the intrinsic $R-L$ scatter is somewhat lower).  \\begin{figure} \\figurenum{7} \\epsscale{1.0} \\plotone{f7.eps}  \\caption{Same as Figure \\ref{Fig:origAcorrmassresid} except the top panel shows the subset of data kept from each sample after removing poor quality data, removing the A11 ``group II'' objects, and using the RM \\Hbeta\\ masses instead of SE masses for the VP06 sample.  The bottom panel shows this same ``high-quality'' sample after applying \\civ\\ mass corrections similar to Equation (1), but determined for each individual sample.  Mean sample offsets are listed but not marked because they are smaller than the typical measurement uncertainty on the mass residuals, $\\sim 0.1-0.2$ dex for the VP06 sample.}  \\label{Fig:bestorigAcorrmassresid} \\end{figure}   \\subsection{The BH Mass -- Luminosity Color Correlation} \\label{S:colorcorrelation}  In their comparison of \\civ\\ and \\Hbeta-based masses in a sample of lensed quasars, A11 discovered a strong correlation between the ratio of \\civ-to-\\Hbeta\\ masses and the ratio of UV-to-optical luminosity.  Such a correlation is naively expected, since the ratio of these masses is a function of the ratio of UV-to-optical luminosity.  However, the slope of the correlation measured by A11 was in excess of that expected, if the only source of the correlation was the dependence of $M_{\\rm BH}$ on $L$ ($M_{\\rm BH} \\propto L^{1/2}$).  Interestingly, if we combine the blueshift--equivalent width relation of \\citet{Richards11} with the blueshift--color relation of \\citet{Gallagher05}, we expect a shape--color correlation in the sense that high (low) $S$ \\civ\\ profiles are seen in bluer (redder) AGNs.  This means that our shape correction operates in the same sense as the A11 color correction.  Figure \\ref{Fig:clrshape} shows this shape--color correlation.  To be consistent, we used only SE, non-host-corrected $\\lambda L_{5100}$ for all objects, which necessarily decreased the VP06 sample by six objects. We compare the statistics from this shape--color correlation with the A11 mass--color correlation, but we first subtracted the expected dependence of the A11 correlated quantities on the ratio of the luminosities by removing the the luminosity term, $(\\lambda L_{1350})^{0.53}/(\\lambda L_{5100})^{0.52}$, from the A11 mass residuals.  This leaves only the excess correlation observed by A11, and is a more consistent comparison with the shape--color correlation, in which there is no redundant dependence on luminosity in both coordinate axes.  The Spearman statistics listed in Table \\ref{Tab_Fits2} show that both correlations are weak, but the correlation of $S$ with color is the stronger of the two.  Since A11 did not use this combined sample to discover or define this correlation, we also make a comparison using the A11 sample alone.  The fit to the A11 sample is denoted by the solid line in Figure \\ref{Fig:clrshape}, and the Spearman statistics are listed the top corner as well as in Table \\ref{Tab_Fits2}.  The statistics show that these correlations are still not particularly strong, but are more significant than for the combined sample (as also noted by A11). The correlation of $S$ with color is, again, the more statistically significant of the two.  \\begin{deluxetable}{llccc} \\tablecolumns{4} \\tablewidth{0pt} \\tablecaption{Spearman Statistics for Correlations with UV-to-Optical Luminosity Ratio} \\tablehead{ \\colhead{Correlated}& \\multicolumn{4}{c}{}\\\\ \\colhead{Quantity}&\\colhead{Data Set}&\\colhead{$r_{\\rm s}$}& \\colhead{$P_{\\rm ran}$} & \\colhead{$N$}}  \\startdata $S$   &   Combined  &  0.231  &  0.152  &  40\\\\ $S$   &   A11       &  0.555  &  0.077  &  11\\\\ $\\log (M_{\\rm CIV}/M_{{\\rm H}\\beta}) -$ &&&& \\\\ \\phn $\\log (L_{\\rm UV}^{0.53}/L_{\\rm opt}^{0.52})$ & Combined & 0.207 & 0.201 & 40\\\\ $\\log (M_{\\rm CIV}/M_{{\\rm H}\\beta}) -$ &&&& \\\\ \\phn $\\log (L_{\\rm UV}^{0.53}/L_{\\rm opt}^{0.52})$ & A11 & 0.527 & 0.096 & 11\\\\ \\enddata   \\label{Tab_Fits2} \\end{deluxetable}   \\begin{figure} \\figurenum{8} \\epsscale{1.0} \\plotone{f8.eps}  \\caption{Comparison between \\civ\\ line shape, $S={\\rm FWHM}/\\sigma_l$, and the ratio of the UV-to-optical luminosities, $L_{1350}/L_{5100}=\\lambda L_\\lambda\\, (1350\\aam )/\\lambda L_\\lambda\\, (5100\\aam )$. Symbol types are the same as in Figure \\ref{Fig:massresidVshape}. The solid line is the best fit to the A11 sample alone.  The statistics in the top left corner are for the A11 sample only; the combined sample has $r_{\\rm s}=0.231$ and $P_{\\rm ran}=0.152$ (see Table \\ref{Tab_Fits2}).}  \\label{Fig:clrshape} \\end{figure}   We also measured the mass--color correlation slope of the A11 Group I objects after correcting these FWHM-based \\civ\\ masses with Equation (1) above.  A11 originally measured this slope to be $0.89 \\pm 0.25$.  We measure a new slope of $0.71 \\pm 0.17$, which is marginally consistent with the expected slope of $\\sim$0.53 from the mass scaling relationship dependence on $L$.  This suggests that the shape bias in the \\civ\\ masses is a likely driver for the A11 mass--color correlation and can at least partially explain the steep slope observed by A11.  Additional work is underway using a larger sample to either identify or rule out additional effects, e.g., non-universal SEDs, that could also be contributing to the excess slope observed by A11.   \\subsection{Origin of the Non-variable C\\,{\\scriptsize IV} Component}  Our analyses in \\S\\S \\ref{S_RMmnrms} and \\ref{S:shapebias} have shown (1) that the SE \\civ\\ profile is a composite of both a variable (reverberating) component and non-variable component, and (2) that object-to-object differences in the characteristics of the non-variable component are a primary cause for the large scatter in FWHM-based SE \\civ\\ BH masses as compared to \\Hbeta\\ masses.  Here we discuss possible origins for these differences.  \\subsubsection{The ``Traditional'' Narrow Line Region} \\label{S_nonBLR}  The first, obvious, possibility is to associate the low velocity component with the ``traditional'', low-density NLR, i.e., the region responsible for emitting narrow forbidden lines, such as \\ob.  A NLR component is observed in the \\Hbeta\\ emission line and has sometimes been assumed to be a component of the \\civ\\ profile as well \\citep[e.g.,][]{Baskin&Laor05, Sulentic07, Greene10, Shen12}.  However, we argue that this cannot be the origin of the non-variable component of \\civ\\ for the following reasons:  \\begin{enumerate} \\item While the C+3 ion is photo-ionized, the \\CIV\\ emission line transition is excited collisionally.  Photoionization models suggest that \\CIV\\ emission from the low-density NLR region must therefore be weak \\citep{Ferland&Osterbrock86}.  In contrast, the non-variable component seen in the mean \\civ\\ profiles of Figure \\ref{Fig:civmnrms} is often quite strong.  Admittedly, \\civ\\ emission is still seen in Seyfert 2 spectra \\citep[e.g.,][]{Collins05, Zheng08}, also suggesting a possible origin in the NLR, but studies have additionally shown that NLR emission line strengths and ratios are different between Seyfert 1's and Seyfert 2's \\citep{Zhang08} and may be dependent on inclination \\citep{Fine11}.  \\item Not all \\civ\\ SE profiles show an obvious narrow, core component. The presence and strength of this component is also anticorrelated with blueshift and the \\civ\\ equivalent width \\citep{Wills93b, Richards02, Leighly04a, Richards11}.  It is much more probable that this reflects differences in the BLR kinematics or geometry between objects rather than correlations between NLR and BLR kinematics.  \\item A large percentage of \\civ\\ SE profile peaks are blueshifted \\citep{Richards11}, some by thousands of km s$^{-1}$.  While NLR emission can be blueshifted from the systemic redshift, it would {\\it not} be expected to show these large of blueshifts.  If present, NLR emission should appear at or at least near the systemic redshift, leaving it relatively redshifted with respect to the highly blueshifted \\civ\\ profile \\citep{Wilkes84}.  One would then expect potential double-peaked profiles in these objects, rather than this component simply being absent.  Such a feature is not observed here, in large samples, or even in composite AGN spectra \\citep[see, e.g.,][]{Richards11}.  \\item We subtracted the scaled rms spectrum from the mean spectrum of all objects shown in Figure \\ref{Fig:civmnrms} (except Fairall 9 because of the poor quality rms spectrum) and measured the FWHM of the residual non-variable components using the same methods as in Section \\ref{S_RMmnrms}.  We also measured the FWHM of the \\oiii\\,$\\lambda$5007 line from this sample, subtracting a linear, local continuum, from available optical data.  Figure \\ref{Fig:nonvarCIVvsOIIIwidths} shows that the velocity widths of the non-variable narrow cores are larger (often significantly so) than those measured in the unblended NLR \\oiii\\,$\\lambda$5007 line. Combining these relatively broad core widths with the core strength and additional blue excess emission in most of this sample, it is unlikely that this component is, at least solely, from the traditional, low-density, extended NLR emission.  \\end{enumerate}  \\begin{figure} \\figurenum{9} \\epsscale{1.0} \\plotone{f9.eps}  \\caption{Comparison between the FWHM of the \\civ\\ non-variable component and the FWHM of the [\\oiii]\\, $\\lambda$5007 line for the RM sample. Individual objects are labeled, and the measurements reflect all objects shown in Figure \\ref{Fig:civmnrms} except Fairall 9, whose rms spectrum was too noisy to accurately isolate the non-variable \\civ\\ component.  The solid line is the 1:1 relation in FWHM, and the two \\civ\\ non-variable FWHM measurements for NGC\\, 4151 connected by the dotted line represent the difference in width if the narrow absorption near 1545\\AA\\ is linearly interpolated across before measuring the FWHM (upper value) or not (lower value).}  \\label{Fig:nonvarCIVvsOIIIwidths} \\end{figure}   \\subsubsection{A BLR Disk-wind} \\label{S_DWBLR}  The presence of a BLR wind has been used to explain a wide variety of observed AGN phenomena.  This includes, for instance, the often blueshifted \\civ\\ profile \\citep[see][and references therein]{Richards11}, absorption line features \\citep[e.g.,][]{Hamann11}, and AGN outflows and feedback \\citep[e.g.,][]{Reeves09}.  \\citet{Richards11} argue that the sources with the largest blueshifted \\civ\\ profiles are the ``wind-dominated'' sources, where this wind component causes a blueshifted peak and line asymmetries \\citep[cf.][]{Leighly04a} in the SE \\civ\\ profile.  These observations are consistent with the non-variable blue excess emission we observe when comparing the RM mean and rms spectra.  However, we also observe a strong low-velocity, non-variable core in some objects.  So in terms of the SE BH mass estimates {\\it alone}, if this total non-variable component (core $+$ blue excess) is a disk-wind, it is the peaky, low $S$-value objects with the strongest SE core components that would have the largest wind ``contamination'' or bias in their FWHM and therefore FWHM-based SE mass measurements.  The BLR disk-wind has been suggested to be launched from the inner, {\\it high-velocity}, high-ionization regions of the BLR or outer accretion disk \\citep[cf.][]{Murray&Chiang97, Elvis00, Elvis12}, so the additional non-variable, {\\it low-velocity} core component observed here does not readily fit into this model unless there is an orientation dependence of the wind component relative to the line of sight (LOS).  The RM mean and rms spectra in Figure \\ref{Fig:civmnrms} as well as \\civ\\ profiles in general show that the strength of the blueward asymmetric emission appears to be roughly anti-correlated with the strength of the low-velocity core (e.g., the blueshift--equivalent width relation discussed by \\citealt{Richards11}).  Therefore, we argue that if this non-variable component does originate in an outflowing BLR wind, there {\\it should} be an orientation dependence, where the change in \\civ\\ profile from `peaky' to `boxy', or low $S$ to high $S$, is due to varying levels of LOS wind contamination to the broad emission line profile as a function of orientation.  Furthermore, because this component does not reverberate, such a wind would necessarily be optically thin to the ionizing continuum \\citep[see, e.g.,][who discuss the possibility that low optical depth can explain the non-variable wings of the \\Hbeta\\ profile]{Korista&Goad04}.  For a disk$+$wind BLR \\citep[cf.\\ Figure 1 of][]{Murray&Chiang97} to support our observations of the variable$+$non-variable \\civ\\ profile and put these observations in context to other \\civ\\ observations, an optically-thin wind that is more polar than equatorial is required.  In this case \\CIV\\ can still be emitted in the wind but does not reverberate.  We can then make a simple generalization for a composite ``disk'' (variable component) plus ``wind'' (non-variable component) \\civ\\ profile \\citep[see also][]{Wang11}.  We would then expect the following orientation-dependence of the profile of each component: \\begin{itemize}  \\item {\\bf The disk:} As the inclination decreases (goes from edge-on to face-on), the disk component experiences a general narrowing in velocity width due to the change in inclination.  There is no other a priori requirement for the shape of the line; photoionization and BLR geometry model-fits to the Balmer line profiles have shown that a wide variety of both single- and double-peaked profiles can exist and still arise from a virialized disk-like distribution of gas with varying optical depths \\citep[see review by][and references therein]{Eracleous09}.  \\item {\\bf The wind:} At high (closer to edge-on) inclinations, the wind component is narrow and centered near the \\civ\\ systemic velocity, as the direction of outflow is largely perpendicular to the LOS.  As the inclination decreases, more of the wind is observed to be outflowing along the LOS.  Consequently, the wind emission is distributed across a larger, preferentially blue-shifted, range of velocities, so the peak also blueshifts and decreases in flux.  \\end{itemize}   \\subsubsection{The Intermediate Line Region}  Another previously suggested explanation for the peaky SE core component in some \\civ\\ lines is that this emission arises in an ``intermediate line region'' \\citep[ILR;][]{Wills93b, Brotherton94}.  Such a region is explained as a higher-velocity, higher-density inner (i.e., closer to the BLR) extension of the NLR.  \\citet{Wills93b} alternately suggest a different, simple model in which the core \\civ\\ emission could be coming from a bipolar outflow, but this may not be mutually exclusive to the ILR/inner NLR.  More recent spatially-resolved studies of the kinematics and physical conditions in the inner NLRs of local Seyferts \\citep[e.g.,][]{Das05, Kraemer09, Crenshaw10, Fischer10, Fischer11} show a link between possible ILR \\civ\\ emission and biconical outflows \\citep[cf.][]{Crenshaw&Kraemer07}.  While we mentioned above that densities in the ``typical'', i.e., extended, NLR are not high enough to result in significant emission from high-ionization species, spatially resolved STIS spectra have co-located high-ionization-line emission with knots of [\\oiii] $\\lambda5007$ emission in the inner NLR \\citep[e.g.][]{Collins05}.  Such emission must be enhanced by collisional processes \\citep{Kriss92}, such as shock heating or microturbulance, mechanisms likely to be present if this emission is arising in an outflow \\citep{Kraemer07}.  \\citet{Nelson00} and \\citet{Kraemer00} show that with {\\it HST} observations of nearby Seyferts, spatially-resolved \\CIV\\ emission does arise outside the unresolved nucleus, i.e., from the ILR/inner NLR.  However, the \\civ\\ emission flux drops significantly faster than \\ob\\ emission as a function of radius.  The spatially resolved inner knots of [\\oiii] $\\lambda5007$ emission also show relatively high, $\\sim$~1000 km s$^{-1}$, and even blueward asymmetric velocities \\citep{Nelson00, Das05, Fischer11}, both reminiscent of the non-variable \\civ\\ component we see here.  Unfortunately, and important for potentially characterizing the non-variable \\civ\\ profile, because the [\\oiii] emission extends to much larger radii and lower velocities, the line profile of the ``integrated'' [\\oiii] emission in typical ground-based spectra of both near and distant objects is not a good template match for the observed non-variable \\civ\\ core emission.  In this type of scenario, the non-variable component of the \\civ\\ line, i.e., both the peaky, low-velocity \\civ\\ core emission and blue excess, could be attributed to the ILR emission and part of a biconical outflow. LOS emission-line profiles from such a distribution of gas would therefore exhibit velocity widths larger than the typical, extended NLR region that is probed by integrated forbidden-line emission profiles, but narrower than typical BLR emission.  Additionally, if the gas is outflowing, blueshifts are to be expected, and potentially more prominent for the higher gas velocities, since this emission likely arises at closer radii and may therefore by more highly collimated by the bicone.  These are all expectations consistent with the observations of the \\civ\\ emission line profile seen here and in previous studies. Furthermore, our orientation-dependent model of the non-variable \\civ\\ emission introduced above fits equally well with an origin in a biconical ILR outflow as in the BLR disk-wind, since the LOS component from the inner region of such an ILR outflow should change as a function of orientation.  Such a distribution of gas would also not be expected to reverberate.  Because the spatial scales are large relative to the time scale of variability, any reverberation signal would be spatially damped.  Therefore, the non-variable nature of such an emission component naturally holds for this origin, at least on reverberation time scales.  \\subsection{Evidence for Orientation-Dependence}  Both origins for the non-variable \\civ\\ component that we suggest above require an orientation dependence to the outflow component of \\civ\\ emission.  We first looked for indications of an orientation dependence by comparing our orientation expectations with the \\civ\\ profile trends observed in the RM sample.  To do this, however, we must isolate the non-variable component profile to search for signs of orientation dependence.  This was done through a spectral decomposition of the mean spectrum \\civ\\ profiles shown in Figure \\ref{Fig:civmnrms}.  We assume that the mean profile is a superposition of the BLR `disk' (variable component) and outflow (non-variable component).  The rms profile shows variable emission only, so we generalize this emission to arise solely in the reverberating BLR `disk'.  We decompose the mean profiles into variable and non-variable components by first fitting the scaled rms profile (solid gray curves in Fig.\\ \\ref{Fig:civmnrms}) with a sixth order Gauss-Hermite polynomial as a means to derive a smooth profile for the disk component.  We subtract this component from the mean spectrum, and attribute the residual non-variable flux to the outflow component. We fit the residual emission again with a GH polynomial after interpolating over any narrow absorption features.  We highlight this decomposition for the \\civ\\ profiles of NGC\\, 4151, whose mean spectrum has a relatively low $S$-value (0.61) with a strong non-variable core, and NGC\\, 7469, which has higher $S$ (0.92) with a notably weaker non-variable core but with additional, non-variable blue excess.  The results of the decomposition for these two sources are shown in Figure \\ref{Fig:diskwindprofile}, where the original mean and rms spectrum and individual and composite (disk plus outflow) profile fits are shown for NGC\\, 4151 and NGC\\, 7469 in the top and bottom panels, respectively.    \\begin{figure} \\figurenum{10} \\epsscale{0.9} \\plotone{f10.eps}  \\caption{Predicted emission line profiles expected if observing an object with a disk-wind BLR at higher, more edge-on inclinations, ({\\it top}) and lower inclinations ({\\it bottom}).  For each panel, the higher flux black curves show the original mean spectrum (also shown in Fig.\\ \\ref{Fig:civmnrms}) used to represent fiducial \\civ\\ profiles for each orientation.  The red curve represents the variable model component fit to the scaled rms spectrum (also in black), the blue curve represents the non-variable model component, and the gray curve is the composite model including both components.  Finally, the dashed curve is the AGN continuum that was subtracted before fitting each emission-line component.}  \\label{Fig:diskwindprofile} \\end{figure}  From this exercise, we observe for NGC\\, 4151 that (1) the rms profile (our disk component) is broad and double-peaked, possibly demonstrating the emission expected from an inclined disk \\citep{Eracleous&Halpern94}, and (2) the relatively narrow velocity distribution of non-variable residuals would also be expected from an outflow directed largely perpendicular to the LOS.  While the model-fit to the rms spectrum over-predicts the flux in the core, this acts only to somewhat underestimate the core strength but does not significantly change the profile of the residual non-variable component.  As it would not be possible to view this broad-line AGN perfectly edge-on and still see a broad-line AGN (in the unified model), the extended and blue-shifted wing in the non-variable component profile is also consistent with the expectations for this component arising in an outflow.  For NGC\\, 7469 we see (1) a less peaky and more blueshifted outflow component and (2) a narrower disk component.  These are both expected characteristics of a more face-on orientation.  The combined fit conspires to produce a smooth but slightly asymmetric profile, characteristic of many SE \\civ\\ profiles.  The orientation-dependent model we describe here is naively simple. Interestingly, however, our crude decomposition using two test cases from the RM sample places no a priori expectations for the profiles of the individual variable and non-variable components.  Yet, the results of attributing these components to a BLR disk$+$outflow composite profile are strikingly similar to what is expected if we are viewing NGC\\, 4151 at a higher inclination (more edge on) and NGC\\, 7469 at a relatively lower inclination (more face on).  There are also numerous discussions in the literature of orientation dependencies to observed AGN properties.  For example, \\citet{Richards02} suggested the possibility of an orientation effect in the \\civ\\ blueshifts, and thus wind component, though interestingly in the opposite sense that we argue here.  \\citet{Richards02} advocated that large blue-shifted profiles were viewed from more edge-on orientations, which follows the expectations of the largely equatorial wind in the \\citet{Murray&Chiang97} model.  However, \\citet{Richards11} argue against such a dependence based on radio morphology studies. Other radio morphology studies also show no or only weak orientation-dependence of \\civ\\ widths \\citep{Jarvis&McLure06, Fine11}. However, these studies are based primarily on using the FWHM, and we have shown here that \\civ\\ FWHM measurements are biased by the non-variable component.  In fact, \\civ\\ profiles resulting from an orientation-dependence of the non-variable component in the sense we propose above (creating preferentially peaky profiles for high inclination and boxy profiles for low inclination) would act to mask the expected changes in the FWHM of the BLR `disk' component that these studies hope to probe:  The non-variable component contamination will lead to an artificially large measurement of the FWHM for low inclinations and an artificially small measurement of the FWHM for high inclinations.  Another place to search for possible orientation-dependent quasar phenomena is in the characteristics of broad absorption line quasars (BALQSOs), which have been inferred to be viewed preferentially more edge-on \\citep[e.g.,][]{Murray95, Elvis00}.  This is partly because their continua are relatively more reddened than non-BALQSOs \\citep[][and references therein]{Reichard03}.  However, other observations and models have shown that (1) polar outflows are also capable of producing BALs \\citep{Brotherton06, Zhou06, Ghosh07, Borguet10}, (2) individual fits to BALs cannot constrain the geometry and orientation of the BAL region \\citep{Hamann93}, and (3) radio BAL and nonBAL quasars are nearly indistinguishable across a wide range of observed and physical properties, including spectral shape, spectral index, and polarization properties \\citep{Bruni12}.  This evidence suggests that BALs can be observed in quasars over broad range of orientations.  Furthermore, discussing the results of some previous studies in the context to this simple orientation dependence of the \\civ\\ SE profile leads to some novel interpretations of the observations (though we do not claim to be able to explain every aspect of these observed phenomena with {\\it only} inclination): \\begin{itemize}  \\item The Baldwin Effect \\citep{Baldwin77}. We argue that the lowest equivalent width objects are those that tend to have high $S$ and are the most face-on.  The continuum photons escaping the central source should therefore be the least obscured by other nuclear material. This will lead to a larger observed luminosity, on average, compared to higher inclination objects, which we argue are observed to have high-equivalent width, low $S$ profiles due to wind contamination.  \\item NLR Component Removal \\citep[e.g.][]{Baskin&Laor05, Sulentic07}. Some studies argue the narrow \\civ\\ core originates in the NLR.  Yet, \\citet{Sulentic07} have pointed out that the narrow-line components in \\civ\\ are often broader and stronger than typical, ``\\oiii\\, $\\lambda$5007-like'', NLR emission.  Though \\citet{Sulentic07} found better correlations between \\civ\\ and \\Hbeta\\ after removing a narrow core component, they were still left with a large scatter between their \\civ\\ and \\Hbeta\\ masses.  In the context of our model and other arguments in Section \\ref{S_nonBLR}, the large scatter remaining in the \\citet{Sulentic07} results is most probably due to the uncertainty in fitting and subtracting their ``narrow'' component.  Since the orientation-dependent non-variable component cannot simply be described as a typical NLR emission line, a traditional spectral decomposition does not work for \\civ.  \\item ``Extincted'' red wing of \\civ\\ \\citep[see e.g., composite spectra of][]{Richards02, Richards11}.  Past physical explanations for this phenomenon \\citep[cf.][]{Gaskell82} include a combination of radial gas motions and obscuration, though Gaskell admits that it is questionable whether the dust grains required for such obscuration could survive in the BLR.  Our orientation-dependent disk-outflow model explains this phenomenon simply: as the inclination decreases, the observed range of disk velocities narrows.  Conversely, the observed velocities of the outflow component are increasing as it becomes directed more along the LOS, effectively filling in the high-velocity portion of the profile previously occupied by the disk emission.  However, this process is preferentially blueward asymmetric because the receding component of the outflow is not visible through the disk.  This leaves the red side of the profile devoid of emission from the highest velocity gas.  \\item Correlation of blueshift with luminosity \\citep{Richards11}. This is expected if these are the population of objects seen more face-on. Furthermore, \\citet{Gallagher05} demonstrated with a color-blueshift distribution that large blueshift AGNs are less likely to have red continua, regardless of whether intrinsic or dust-reddened.  This would similarly be true if we are seeing these objects more face-on.  \\item Intrinsic colors of BALQSOs \\citep{Reichard03}.  Despite the observed trend of BALQSOs being more reddened than nonBAL quasars, \\citet{Reichard03} found that after correcting for SMC-like dust reddening, BALQSOs appear to be drawn from the same parent population as nonBALs \\citep[see also][]{Bruni12}.  Upon further detailed inspection of BALQSO properties as a function of `intrinsic' colors, \\citet{Reichard03} also discovered the same nonBAL quasar \\civ\\ emission line trends of larger (smaller) equivalent width profiles for `intrinsically' red (blue) HiBALs (their Figure 9).  This suggests to us that the obscuration typical of BALQSOs may not necessarily be an orientation dependent effect (e.g., from a torus, which would indicate a preferential edge-on orientation).  Instead, if the \\civ\\ line shape trend {\\it is} an orientation effect as we suggest, the obscuring regions responsible for producing BALs in quasars may in fact be orientation independent. The range of observed \\civ\\ profiles in BALQSOs is then additional evidence that BALs can be observed over a range in orientations.  \\end{itemize}  Realistically, there are many geometrical, kinematical, and physical properties of the BLR that could cause varying contributions from the LOS component of a BLR disk-wind, and most likely, more than one effect is present in varying degrees for different samples of AGNs.  For example, \\citet{Murray&Chiang97} find that stronger core emission from an equatorial wind component can occur without requiring an inclination-dependence; instead, it is the result of simply extending the outer radius of the wind-emitting region.  We focus only on inclination here because it is arguably the simplest to conceptualize, still suffices to explain many observed trends, and fits the observations whether the non-variable emission originates in an outflow from the BLR or ILR.  Ideally, we would like to use direct orientation measurements to corroborate our interpretation that the \\civ\\ profile changes are a result of an orientation-dependence of the non-variable emission.  Unfortunately, such orientation measurements are scarce.  The most promising methods use radio morphology \\citep[e.g.][]{Brotherton96, Vestergaard00, Jarvis&McLure06, Fine11} or NLR kinematics \\citep[see, e.g.,][]{Fischer10, Fischer11}.  As mentioned above, past radio studies may be biased by typically only using the FWHM, and these studies should be revisited \\citep[though see][who find the strongest \\civ\\ orientation dependence using the full width at $20\\%-30\\%$ peak intensity, which would be less susceptible to contamination from the non-variable core]{Vestergaard00}.  NGC\\, 3783 and NGC\\, 4151, two objects in our RM sample, have been targeted for NLR kinematical modeling to determine orientation.  The best-fit models indicate an inclination of only $15^{\\circ}$ away from the LOS for NGC\\, 3783 (T.\\ Fischer, priv.\\ comm.; Fischer et al.\\ in prep) and $45^{\\circ}$ for NGC\\, 4151, near the maximum expected for a Type 1 spectrum \\citep{Das05}.  Putting the relative strength of the non-variable core observed in each of these objects in Figure \\ref{Fig:civmnrms} in context to these independent orientation estimates seems to support our expectations that we are observing NGC\\, 3783 and NGC\\, 4151 at relatively low and high inclination, respectively."
    },
    "1208/1208.3523_arXiv.txt": {
        "abstract": "In the companion Paper~XV of this series we derive accurate total mass-to-light ratios $(M/L)_{\\rm JAM}\\approx(M/L)(r=\\re)$ within a sphere of radius $r=\\re$ centred on the galaxy, as well as stellar $(M/L)_{\\rm stars}$ (with the dark matter removed) for the volume-limited and nearly mass selected (stellar mass $M_\\star\\ga6\\times10^9 \\msun$) \\atl\\ sample of 260 early-type galaxies (ETGs, ellipticals Es and lenticulars S0s). Here we use those parameters to study the two orthogonal projections $(\\mjam,\\se)$ and $(\\mjam,\\rmaj)$ of the thin Mass Plane (MP) $(\\mjam,\\se,\\rmaj)$ which describes the distribution of the galaxy population, where $\\mjam\\equiv L\\times (M/L)_{\\rm JAM}\\approx  M_\\star$. The distribution of galaxy properties on both projections of the MP is characterized by: (i) the same zone of exclusion (ZOE), which can be transformed from one projection to the other using the scalar virial equation. The ZOE is roughly described by two power-laws, joined by a break at a characteristic mass $\\mjam\\approx3\\times10^{10} \\msun$, which corresponds to the minimum \\re\\ and maximum stellar density. This results in a break in the mean $\\mjam-\\se$ relation with trends $\\mjam\\propto\\sigma_{\\rm e}^{2.3}$ and $\\mjam\\propto\\sigma_{\\rm e}^{4.7}$ at small and large \\se\\ respectively; (ii) a characteristic mass $\\mjam\\approx2\\times10^{11} \\msun$ which separates a population dominated by flat fast rotator with disks and spiral galaxies at lower masses, from one dominated by quite round slow rotators at larger masses; (iii) below that mass the distribution of ETGs properties on the two projections of the MP tends to be  constant along lines of roughly constant $\\sigma_{\\rm e}$, or equivalently along lines with $\\rmaj\\propto\\mjam$ respectively (or even better paralell to the ZOE: $\\rmaj\\propto M_{\\rm JAM}^{0.75}$); (iv) it forms a continuous and parallel sequence with the distribution of spiral galaxies; (v) at even lower masses, the distribution of fast rotator ETGs and late spirals naturally extends to that of dwarf ETGs (Sph) and dwarf irregulars (Im) respectively.  We use dynamical models to analyse our kinematic maps. We show that \\se\\ traces the bulge fraction, which appears the main driver for the observed trends in the dynamical $(M/L)_{\\rm JAM}$ and in indicators of the $(M/L)_{\\rm pop}$ of the stellar population like H$\\beta$ and colour, as well as in the molecular gas fraction. A similar variation along contours of \\se\\ is also observed for the mass normalization of the stellar Initial Mass Function (IMF), which was recently shown to vary systematically within the ETGs galaxy population. Our preferred relation has the form  $\\log_{10} [(M/L)_{\\rm stars}/(M/L)_{\\rm Salp}]=a+b\\times\\log_{10}(\\se/130\\, \\kms)$ with $a=-0.12\\pm0.01$ and $b=0.35\\pm0.06$. Unless there are major flaws in all stellar population models, this trend implies a transition of the mean IMF from Kroupa to Salpeter in the interval $\\log_{10}(\\se/\\kms)\\approx1.9-2.5$ (or $\\se\\approx90-290$ \\kms), with a smooth variation in between, consistently with what was shown in \\citet{Cappellari2012}. The observed distribution of galaxy properties on the MP provides a clean and novel view for a number of previously reported trends, which constitute special two-dimensional projections of the more general four-dimensional parameters trends on the MP. We interpret it as due to a combination of two main effects: (i) an increase of the bulge fraction, which increases \\se\\ and decrease in \\re, and greatly enhances the likelihood for a galaxy to have its star formation quenched, and (ii) dry merging, increasing galaxy mass and \\re\\ by moving galaxies along lines of roughly constant \\se\\ (or steeper), while leaving the population nearly unchanged. ",
        "introduction": "Much of our understanding of galaxy formation and evolution comes from the study of dynamical scaling relations relating galaxy luminosity or mass, size and kinematic \\citep[e.g.][]{Faber1976,Kormendy1977,Dressler1987,Faber1987,Djorgovski1987} or regular trends in the distribution of galaxy properties as a function of their scaling parameters \\citep[e.g.][]{Bender1992,Burstein1997,Kauffmann2003sfh,Gallazzi2006}, and from the study of their evolution with redshift \\citep[e.g.][]{vanDokkum1996,Kelson1997,vanDokkum1998,Treu2005,Franx2008}.  The volume-limited \\atl\\ sample of nearby early-type galaxies (\\citealt{Cappellari2011a}, hereafter Paper~I) constitute an ideal benchmark for studying scaling relations and the distribution of galaxy properties, given the  availability of a high-quality multi-wavelength dataset. In particular in \\citet[hereafter Paper~XV]{Cappellari2012p15} we used the state-of-the-art integral-field kinematics, in combination with detailed axisymmetric dynamical models, to derive accurate masses and global dynamical parameter. We found that galaxies lie, with very good accuracy, on a thin Mass Plane (MP) describing galaxies in the parameter space defined by mass, velocity dispersion and projected half-light radius $(\\mjam,\\se,\\rmaj)$. Here \\mjam\\ is a very good estimate of the galaxy total stellar mass, \\rmaj\\ is the major axis of the `effective' isophote containing half of the observed galaxy light and \\se\\ is the velocity dispersion measured measured within that isophote. The existence of this MP is mainly due to the virial equilibrium condition \\begin{equation}\\label{eq:virial} M_{\\rm JAM}\\propto\\sigma_{\\rm e}^2 R_{\\rm e}^{\\rm maj} \\end{equation} combined with a smooth variation of galaxy properties with \\se. For this reason by itself it contains no useful information on galaxy formation. All the useful constraints on galaxy formation models come from the inhomogeneous distribution of galaxies in non-edge-on views of the MP and from the distribution of galaxy properties along the MP.  This paper is devoted to a study of the non edge-on projections of the MP to see what we learn from it on galaxy formation. This is done in the spirit of the classic papers by \\citet{Bender1992} and \\citet{Burstein1997}. However the fact that we have accurate dynamical masses implies that our MP is extremely thin and it follows the scalar virial relation~(\\ref{eq:virial}) quite accurately. For this reason we can ignore edge-on views of the plane and focus on non edge-on projections only. The thinness of the MP implies that any inclined projection shows essentially the same information, after a change of coordinates. We can use standard and easy-to-understand observables as our main coordinates, instead of trying to observe the plane at a precisely face-on view.  In this paper, in Section~2 we summarize the sample and data, in Section~3 we present our projections of the MP. We illustrate the distribution of a number of quantities on the $(\\mjam,\\se)$ and $(\\mjam,\\rmaj)$ projection. We show the variation of the $(M/L)_{\\rm JAM}$, as well as of population indicators of $M/L$. We show the variation of galaxy concentration, intrinsic shape, morphology and stellar rotation. The variations of the IMF are separately presented in Section~4, together with a review of previous results on the IMF variation. in Section~5 we discuss the implications of our findings for galaxy formation and briefly summarize our paper in Section~6. ",
        "conclusions": "\\subsection{Previous relations as seen on the Mass Plane}   We have shown in Paper~XV that the galaxy Fundamental Plane \\citep{Dressler1987,Djorgovski1987} can be accurately explained by virial equilibrium combined with a smooth variation of galaxy properties, mainly the total mass-to-light ratio $M/L$, with velocity dispersion, with galaxies lying on a tight MP $(M_{\\rm JAM},\\sigma_{\\rm e},R_{\\rm e}^{\\rm maj})$, for a large volume-limited sample of ETGs (Paper~I). Once this has been established, the interesting information on galaxy formation is then all contained in the distribution and the physical properties of galaxies on this plane, which we presented in this paper.  We find that on the MP: (i) galaxies sizes are delimited by a lower-boundary, which has a minimum at a characteristic mass $M_{\\rm JAM}\\approx3\\times10^{10} \\msun$; (ii) A number of key galaxy properties: dynamical $M/L$, and its population indicators H$\\beta$ and colour, as well as the molecular gas fraction, which are mainly related to age, the normalization of the IMF and the prominence of the bulge, all tend to be constant along lines of constant $\\sigma_{\\rm e}$, on the MP; (iii) Another characteristic mass for galaxy properties is the value $M_{\\rm JAM}\\approx2\\times10^{11} \\msun$ which separates a region dominated by the round or weakly triaxial slow rotators at large masses from one dominated dominated by fast-rotators ETGs, flattened in their outer parts and with embedded exponential disks (Paper~XVII), whose characteristics merge smoothly with the ones of spiral galaxies. A transition in the bulge fraction at this galaxy mass appears required in our models for the formation of fast and slow rotators \\citep[hereafter Paper VIII]{Khochfar2011}.  Although our \\atl sample is limited to a minimum mass $M_{\\rm JAM}\\ga 6\\times10^9$ \\msun, our picture naturally extends to lower masses. As shown in \\reffig{fig:mass_size_spirals_etgs} our trends for fast rotator ETGs continues with the dwarf spheroidal (Sph) sequence at lower masses (see also fig.~7 of \\citealt{Binggeli1984}; fig.~38 of \\citealt{Kormendy2009}; fig.~12 of \\citealt{Chen2010}; fig.~4 of \\citealt{Misgeld2011}; fig.~20 of \\citealt{Kormendy2012}) , while the spirals sequence continues with a sequence of low-mass late spirals or irregulars (Sc--Irr), as independently noted also by \\citet{Kormendy2012}. Interestingly the approximate stellar mass $M_\\star\\approx2\\times10^9$ \\msun\\ where there is a break in the M-size relation of dwarf galaxies and where galaxies with bulge starts to appear (\\reffig{fig:mass_size_spirals_etgs}) corresponds to the threshold for quenching of field galaxies recently discovered by \\citet{Geha2012}. Below that mass only the cluster or group environment can strip spirals of their gas. Perhaps below that mass bulges cannot grow and star formation cannot be quenched by internal processes.  \\begin{figure*} \\includegraphics[width=0.8\\textwidth]{fig14} \\caption{Schematic summary of the results presented in \\refsec{sec:vp_projections} and \\refsec{sec:imf}. ETGs properties, dynamical $M/L$ (\\reffig{fig:virial_plane_projections_ml}) or its population proxies, H$\\beta$ and galaxy colour, as well as the molecular gas fraction (\\reffig{fig:virial_plane_projections_hbeta}), kinematical concentration (\\reffig{fig:virial_plane_projections_sig8_sige}), which traces bulge mass, and IMF mass normalization (\\reffig{fig:vp_projections_imf}), all tend to vary along lines with roughly $\\re\\propto M$ (or even better $\\re\\propto M^{0.75}$), where \\se\\ is nearly constant. This sequence of ETGs properties merges smoothly with the one of spiral galaxies, with little overlap between late spirals (Sc-Irr) and ETGs, a significant overlap between early spirals (Sa-Sb) and fast-rotator ETGs with low $M/L$ and no overlaps between spirals and fast-rotators with high $M/L$. Three characteristic masses are emphasized in this diagram: (i) below $M_\\star\\approx2\\times10^9$ \\msun\\ there are no regular ETGs and the mass-size lower boundary is increasing; (ii) $M_\\star\\approx3\\times10^{10}$ \\msun\\ is the mass at which ETGs reach their minimum size (or maximum stellar density), before a sudden change in slope $\\re\\propto M^{0.75}$ at larger masses (see also fig.~4 in Paper~I); (iii) Below $M_\\star\\approx2\\times10^{11}$ \\msun\\ ETGs are dominated by flat fast rotators, showing evidence for disks (Paper~XVII), while slow rotators are rare. Above this mass there are no spirals and the population is dominated by quite round or weakly triaxial slow rotators (paper~III) with flat (core/deficit) central surface brightness profiles (Paper~XXIV). These smooth trends in scaling relations motivated our proposed parallelism between spirals and ETGs. To emphasize this connection we uses the same morphology symbols here as in our `comb' diagram in fig.~2 of Paper~VII.} \\label{fig:etgs_spirals_diagram} \\end{figure*}  The parallelism between fast rotator ETGs and spiral galaxies in scaling relations was one of the driver for our proposed revision (Paper~VII) to the tuning-fork diagram by  \\citet{Hubble1936}, in the spirit of \\citet{vandenBergh1976} proposed parallelism between S0s and spirals. In \\reffig{fig:etgs_spirals_diagram}. we summarize our results on morphology, kinematics, population and scaling relations in a single diagram, using the same galaxy symbols as fig.~2 of Paper~VII, to provide a link between the two papers. This picture allows us to provide a new perspective and a clean empirical view of a number of classic scaling relation and known trends in galaxy properties.  \\subsubsection{$L-\\sigma$ trends}  The classic \\citet{Faber1976} relation $L\\propto\\sigma^4$ is well known to be a projection of the FP. We study it here using dynamical mass instead of light. Our study for the first time uses \\se\\ values integrated within \\re, which properly account for both velocity dispersion and stellar rotation, in combination with accurate dynamical masses which, unlike luminosity, can be properly related to the stellar kinematics. We find that the relation is not well described by an approximately linear trend, but it shows a break which is a reflection of the one in the ZOE (equation \\ref{eq:zoe}). The slope of the mean $\\sigma-M$ relation (equation \\ref{eq:m-sigma}) changes from $\\mjam\\propto\\sigma_{\\rm e}^{2.3}$ below $\\mjam\\approx5\\times10^{10}$ \\msun\\ ($\\se\\approx140$ \\kms) to $\\mjam\\propto\\sigma_{\\rm e}^{4.7}$ at larger masses. Thanks to our use of dynamical masses and \\se\\ we find for the first time that the bend in the relation is quantitatively consistent with the one we observe in the M-size relation.  The slope we find at the high-mass range is consistent with the original \\citet{Faber1976} relation, while the low-mass slope in our $\\sigma-M$ relation is consistent with claims of a change in the slope of the $\\sigma-L$ relation of elliptical galaxies at low luminosities, where the relation was reported to become $L\\propto\\sigma^2$ \\citep{Davies1983,Held1992,Matkovic2005,deRijcke2005,Lauer2007,Forbes2008,Tortora2009}. However previous studies differ from ours because of $M/L$ variations, they ignore stellar rotation, suffer from small number statistics and selection biases. This may be the reason why previous studies suggested a value $M_B\\approx-20.5$ mag for the break in the $\\sigma-L$ relation \\citep[see Sec.~3.3.3 of][for a review]{Graham2011}, which corresponds to the characteristic mass $M\\approx2\\times10^{11}$ \\msun, using the $(M/L)-\\sigma$ relation \\citep{Cappellari2006} and considering a typical galaxy colour $B-I\\approx2.2$ mag. The previously reported mass is clearly inconsistent with our accurate value $M_{\\rm JAM,b}\\approx5\\times10^{10}$ \\msun. Our new result shows that the break in the relation is associated to the characteristic mass $\\mjam\\approx3\\times10^{10}$ \\msun\\ of the break in the ZOE and not with the other characteristic mass $\\mjam\\approx2\\times10^{11}$ \\msun\\ above which round, slow rotator, cored galaxies dominate.  Observations from the much larger SDSS sample have failed to find evidence for a clear break in the $\\sigma-L$ relation \\citep{Bernardi2003fp,Gallazzi2006,Hyde2009curv}. For this reason the relation is still often assumed to be a power-law. The accuracy and homogeneity of our data, the size of our sample and especially the  quantitative consistency between the $\\sigma-M$ and $\\re-M$ relation provides an unambiguous confirmation for the break and an accurate determination of the transition mass.  \\subsubsection{$L-\\re$ trends}  Another well known projection of the FP is the \\citet{Kormendy1977} relation. When using mass instead of light it becomes clear it represents the analogue of the \\citet{Faber1976}, but this time in the $(M_{\\rm JAM},R_{\\rm e}^{\\rm maj})$ projection of the MP. Also in this case, when samples are morphologically selected to consist of ellipticals, they tend to populate mostly the region of the diagram near the ZOE, defining a relatively narrow sequence \\citep{Graham2008curv,Kormendy2009,Chen2010,Misgeld2011}. Although the sequence is useful for a number of studies, it is important to realize that it is not a real sequence in galaxy space on the MP. It is due to the sample selection and it represents essentially one of the contour levels of a continuous trend of galaxy properties, spanning from spiral galaxies, to ETGs, and only terminating on the well defined ZOE (\\reffig{fig:mass_size_spirals_etgs} here and figure~4 in Paper~I).  Given that photometry is much easier to obtain than stellar kinematics, a change of slope in the luminosity-size relation has been known for long. It was pointed out by \\citet[their fig.~7]{Binggeli1984} when combining photometry measurements of dwarf spheroidals and ordinary ellipticals: the dwarf spheroidals sequence appears to sharply bend from the ellipticals sequence. Similar differences in the slope of the luminosity-size relation of dwarfs and ellipticals were presented by a number of authors \\citep[e.g.][]{Kormendy1985,Graham2003,Kormendy2009,Misgeld2011}. The change of slope has been interpreted in different ways.  \\citet{Kormendy1985} and \\citet{Kormendy2009} interpreted dwarf spheroidal as constituting a separate family, of gas-stripped dwarf spirals/irregulars, while \\citet{Graham2003} and \\citet{Graham2008curv} explain the change of slope or curvature in the relation as a natural consequence of the variation of the \\citet{Sersic1968} index with luminosity \\citep[e.g.][]{Young1994,Graham2003} in a homogeneous class of elliptical galaxies with a range of masses (see \\citealt{Graham2011} and \\citealt{Kormendy2012} for two complementary reviews of this subject).  Our results cannot be compared to theirs in a statistical sense, as galaxies in their diagrams are, by design, not representative of the population in the nearby Universe, and certain classes of objects (e.g. M32) are over-represented. Our sample is volume-limited and for this reason it gives a statistically representative view of the galaxy population above a certain mass. Still the fact that the sequence of dwarf spheroidals and low-mass spirals/irregulars, lie on the continuation of our trends for fast rotator ETGs and spiral galaxies respectively (\\reffig{fig:mass_size_spirals_etgs}), below the $M_\\star\\la6\\times10^{9} \\msun$ mass limit of our survey, suggests a continuity between dwarf spheroidals and the low-mass end of our disk-dominated fast-rotator ETGs population, which in turns we showed are closely related to spiral galaxies. For this reason our results reconciles the apparent contrast between the findings of a Sph-E dichotomy \\citep{Kormendy1985,deRijcke2005,Janz2008,Kormendy2009} and the ones of a continuity \\citep{Graham2003,Gavazzi2005,Ferrarese2006acs,Cote2006,Janz2009,Boselli2008,Graham2008curv,Forbes2011}. Our finding in fact agrees with the proposed common origin of dwarf spheroidal and low-mass spiral galaxies and irregulars \\citep{Kormendy1985,Dekel1986}, but also shows an empirical continuity between dwarf spheroidals and a subset of the ellipticals family. The missing link between Sph and E is constituted by disk-dominated fast rotator ETGs \\citep{Emsellem2007,Cappellari2007}. In fact the continuity we find is not between ``true'' ellipticals, namely the slow rotator, and dwarf spheroidals, but between ``misclassified'' ellipticals with disks and S0, namely the fast rotators, and dwarf spheroidals. After this text was written and the parallelism between spiral galaxies and early-type galaxies were presented (Paper~VII) to interpret the trends we observed in galaxy scaling relations (Paper~I and \\citealt{Cappellari2011dur}), an independent confirmation of this picture, including its extension to low mass systems was also provided by \\citet{Kormendy2012}.  \\subsubsection{Population trends with $\\sigma$ or \\re}  The characteristic mass $M_{\\rm JAM}\\approx3\\times10^{10} \\msun$ is the same transition mass discovered by \\citet{Kauffmann2003mass} who state that ``low-redshift galaxies divide into two distinct families at a stellar mass of $3\\times10^{10} \\msun$. Lower-mass galaxies have young stellar populations, low surface mass densities and the low concentrations typical of discs. Their star formation histories are more strongly correlated with surface mass density than with stellar mass.'' A similar trend involving colours and also better correlated with surface density (or with the velocity dispersion estimated from the photometry) than with mass, was found to extend to redshift up to $z\\approx3$, with red galaxies being systematically small, and blue galaxies being large at a given mass \\citep{Franx2008}. This was recently confirmed, still using photometric data alone, by \\citet{Bell2012}. The correctness of all these statements can now be easily and accurately verified for the nearby ETGs subset in \\reffig{fig:virial_plane_projections_ml} and \\ref{fig:virial_plane_projections_hbeta}.  The novelty of our work, with respect to all previous studies, is that we have unprecedentedly accurate and unbiased dynamical masses, and stellar velocity dispersions \\se\\ integrated within a large aperture (1\\re), instead of $\\sigma$ values inferred from photometry. This allow us to conclusively state that neither dynamical mass nor stellar surface density are actually the best descriptor of galaxy properties, the main trend being along the $\\sigma_{\\rm e}$ direction (which includes rotation in the case of disk galaxies). Our clean \\atl\\ result was already presented in \\citet{Cappellari2011dur} and subsequently confirmed with SDSS data and using virial mass estimates by \\citet{Wake2012}. The trend we observed for the ETGs can be extended to spiral galaxies, which fill the region of larger sizes above the ETGs in the $(M_{\\rm JAM},R_{\\rm e}^{\\rm maj})$ projection, smoothly overlapping with the ETGs for the largest spiral bulge fractions (figure~4 of Paper~I and \\reffig{fig:mass_size_spirals_etgs}). It has been known for long that in spirals luminosity-weighted ages are lower, star formation is larger and colours bluer on average than the ETGs, essentially by definition \\citep[e.g.][]{Hubble1936,vandenBergh1976}.  Similar trends between age and size (or surface brightness), with older objects being smaller at given age, were found to persist in ETGs. \\citet{vanderWel2009} state that ``at a given stellar velocity dispersion, SDSS data show that there is no relation between size and age''. The same age-size trend was pointed out by \\citet{Shankar2009}, and in different terms by \\citet{Graves2009b} who state that ``no stellar population property shows any dependence on \\re\\ at fixed $\\sigma$, suggesting that $\\sigma$ and not dynamical mass is the better predictor of past SFH''. These findings are another way of saying that age variations must follow lines of constant $\\sigma$ on the MP as we find here for $M/L$, H$\\beta$, colour and molecular gas fraction, and confirm using age for our sample in McDermid et al. (in preparation), in agreement with \\citet[their figure~15]{Gallazzi2006}.  The age-size trend was also confirmed in different samples of nearby galaxies by \\citet{Valentinuzzi2010} and \\citet{Napolitano2010}, and a similar trend, with star forming galaxies being larger than passive ones was found in place from redshift $z\\sim2.5$ \\citep{Williams2010z2,Wuyts2011,Newman2012}. The only contrasting view is the one by \\citet{Trujillo2011}, who find a lack of age-size trend both at low and high-redshift.  The same characteristics mass $M_\\star\\approx3\\times10^{10}$ \\msun\\ of \\citet{Kauffmann2003mass}, which constitute the location of the break in our ZOE, was found by \\citet{Hyde2009curv} in the mass-size relation of $5\\times10^4$ SDSS galaxies. Their trend is significant but quite subtle. The reason for the curvature they find becomes clear from what we find: the average radius of the objects on the MP at constant mass, bends upwards at low masses, due to the cusp in the ZOE, with the strength of the effect being dependent on the specific criterion adopted to select ETGs.  \\subsection{Implications for galaxy formation}  Galaxy formation is the superposition of a number of complex events that happen in parallel. Here we sketch a tentative picture of some of the phenomena that can play a role in explaining what we see.  We refer the reader to \\citet[hereafter Paper~VI]{Bois2011}, Paper~VIII and Naab et al. (in preparation) for a more in-depth discussion.  The smoothness and regularity of the trends we observe and the fact that they extend to spirals, appears to indicate a close connection between the formation of the two classes of objects  (Paper~VIII). The same similarity between the fast rotator ETGs and spirals in terms of their morphology and degree of rotation, lead us to propose a revision (Paper~VII) of the classic morphological classification \\citep{Hubble1936} to emphasize the parallelism between the fast rotators and spirals, in the same way that \\citet{vandenBergh1976} proposed it for S0s and spirals. Only the most massive slow rotators appear to form an empirically separated class. The {\\em kinematic} morphology-density relation (Paper~VII), which applies to our kinematic classes the relation discovered by \\citet{Dressler1980}, suggest that most spirals are being transformed into fast rotators due to environmental effects \\citep{Khochfar2008}, with a mechanism that is sufficiently `gentle' to preserve the near axisymmetry of the disk (Paper~II).  One process which is often mentioned in the context of cluster galaxy populations is stripping by the interstellar medium \\citep{Spitzer1951,Abadi1999}, which may impact the global morphology of the galaxy by removing a significant fraction (or most) of its gas content but mostly preserving the stellar disk component. This may explain some of the most flattened galaxies in the fast rotator class (e.g., NGC4762), as emphasised in Paper~VII. This would work by removing some of the mass of the galaxy keeping a relatively constant effective size (of the stellar component), and may contribute to the scatter in the Mass-size plane for fast rotator as well as to their overlap with the more gas-rich spiral galaxies.  As to explain the intermediate to high mass end of the fast rotator class, we would need a process which is able to increase the bulge size, while at the same time removing the gas or shutting off star formation (Paper~VIII). The empirical signatures of this phenomenon are visible in our data as an 'apparent' decrease of the galaxy size, and increase of $\\sigma_{\\rm e}$, which is actually due to the bulge concentrating more light at smaller radii, accompanied by a increase in $M/L$, which is tracing a decrease of star formation or an age increase. The process appears to generally preserve the intrinsic flatness of the stellar disks at large radii (top panel of \\reffig{fig:virial_plane_projections_concentration}). A similar scenario was recently proposed by \\citet{Bell2012} to interpret the relationship between rest-frame optical colour, stellar mass, star formation activity, and galaxy structure from $z\\approx2$ to the present day.  Intense gas-rich accretion events, mostly via cold streams \\citep{Keres2005,Dekel2009}, or major gas rich mergers (Paper~VIII), will increase both the mass and $\\sigma_{\\rm e}$. During the accretion the gas may sink toward the centre \\citep{Mihos1994burst} until it becomes self-gravitating and starts forming stars. It is during this phase of rapid gas accretion that the $(M/L)-\\sigma$ \\citep{Cappellari2006,vanderMarel2007} and ${\\rm IMF}-\\sigma$ (\\refsec{sec:imf_sigma}) relations, the tilt of the FP \\citep{Dressler1987,Faber1987,Djorgovski1987}, the ${\\rm Mg}b-\\sigma$ \\citep{Burstein1988pop,Bender1993pop} or the ${\\rm Mg}b-V_{\\rm esc}$ relation (\\citealt{Davies1993}; \\citealt{Scott2009}, Paper~XXI), will be imprinted in the ETGs population \\citep{Robertson2006fp,Hopkins2009extralight} and then mostly preserved in the following evolution.  The early progenitors of today's fast rotator ETGs would be high-redshift spirals, which are different from local ones. They are characterized by giant gas clumps \\citep{Elmegreen2007,Genzel2011} have high gas fractions \\citep{Tacconi2010,Daddi2010}, possess large velocity dispersion and are dominated by turbulent motions \\citep{ForsterSchreiber2006,ForsterSchreiber2009,Genzel2006,Genzel2008,Law2012}. In that situation bulges may form naturally as the clumps collide and sink to the centre \\citep{Bournaud2007clumps,Dekel2009clumps}, unless they are efficiently destroyed by stellar feedback \\citep{Genel2012,Hopkins2011}. Secular effects \\citep{Kormendy2004} will also contribute to the bulge growth and $\\sigma_{\\rm e}$ increase, while keeping the mass unchanged, as will contribute the destabilizing effect of minor mergers.  \\begin{figure} \\plotone{fig15} \\caption{Evolution scenario for ETGs. The symbols are the same as in \\reffig{fig:virial_plane_projections_ml}, while the large arrows indicate the proposed interpretation of the observed distribution as due to a combination of two processes: (a) in-situ star formation: bulge or spheroid growth, which seems associated to the quenching of star formation, which moves galaxies to the right of towards the bottom, due to the increased concentration (decreasing \\re\\ and increasing \\se), and (b) external accretion: dry mostly minor merging, increasing \\re\\ by moving galaxies along lines of roughly constant \\se\\ (or steeper), while leaving the population unchanged. A schematic illustration of these two processes is shown in \\reffig{fig:black_holes_growth}.} \\label{fig:etgs_formation_scenario} \\end{figure}  \\begin{figure} \\plotone{fig16} \\caption{Schematic representation of the two main processes responsible for the formation of the observed distribution of galaxies on the MP. (a) In-situ star formation: bulge growth via cold accretion, secular evolution, or minor gas-rich mergers, followed by quenching by AGN or other mechanisms, leaving the galaxy more massive, more compact, and consequently with a larger $\\sigma_{\\rm e}$, and gas poor (blue arrow in \\reffig{fig:etgs_formation_scenario}); (b) external accretion: major or minor dry mergers, increasing galaxy mass and sizes at nearly constant $\\sigma_{\\rm e}$, or with a possible decrease, leaving the population mostly unchanged (red arrow in \\reffig{fig:etgs_formation_scenario}). (taken from \\citealt{Cappellari2011nat}).} \\label{fig:black_holes_growth} \\end{figure}  During the bulge growth some process must be able to turn off star formation, without destroying the fast rotating disks that still dominate the local ETGs population (Paper~II; Paper~III) and that dominates the ETGs population already from $z\\sim2$ (\\citealt{vanderWel2011}; see also \\citealt{vanDokkum2011nat}). The reduced efficiency of star formation by morphological quenching \\citep{Martig2009,Ceverino2010} may be one of the processes explaining why on average bigger bulges correspond to older ages and larger $M/L$. However we showed in \\reffig{fig:virial_plane_projections_hbeta} that the cold gas fraction is decreasing with the bulge fraction, which shows that less fuel is available for galaxies with more massive bulges. Bulge/spheroid growth seems also associated with AGN feedback, which would provide another bulge-related quenching mechanism \\citep{Silk1998,Granato2004,Bower2006,Hopkins2006,Croton2006}. During the sequence of bulge/spheroid-growth followed by quenching, the original gas rich spiral will move from the left of the observed  $(M_{\\rm JAM},R_{\\rm e}^{\\rm maj})$ plane towards the right, or possibly the bottom right due to the increased concentration, while intersecting the constant $\\sigma$ lines (blue arrow in \\reffig{fig:etgs_formation_scenario}). At the end of the evens the galaxy will be a fast rotator ETG, generally more massive and with a bigger bulge (smaller \\re, larger \\se\\ and concentration) than the precursor clumpy spiral.  At any stage during the bulge-growth phase the galaxy may accrete purely stellar satellites \\citep{Khochfar2006b}. In the case of these dry mergers the situation is quite different and one can predict the final configuration using energy conservation \\citep{Hernquist1993,Cole2000,Boylan-Kolchin2006,Ciotti2007}. The predictions show that sizes will increase as the mass grows. For major  mergers (equal mass) and typical orbital configurations, one can show that the mass and radius double, while $\\sigma$ remains constant in agreement with simulations (\\citealt{Nipoti2009}; see \\citealt{Hilz2012a} and \\citealt{Hilz2012b} for more accurate numerical simulations and detailed physical explanations of this process). While in the limit in which the same mass doubling happens via small satellites, as mostly expected from the shape of the \\citet{Schechter1976} mass function, the radius will increase by a factor of four and the dispersion will be twice smaller \\citep{Naab2009,Bezanson2009,Hopkins2009compact,Hilz2012a}. As a result the galaxy will move along lines that are parallel to the constant $\\sigma$ lines, or steeper. During these dry mergers, given that there is little gas involved, the stellar population, colour and $M/L$ will remain unchanged (red arrow in \\reffig{fig:etgs_formation_scenario}).  When the galaxy grows more massive than the characteristic mass $M_\\star\\ga2\\times10^{11}$ \\msun, a transition seems to happen in their formation, possibly associated to the shock heating of the infalling gas, due to the halo mass \\citep{Dekel2006}. Above this mass galaxies are in fact embedded in massive and hot X-ray halos (\\citealt{Kormendy2009}, Paper~XIX) which prevents any further cold gas accretion onto the galaxy.  The gas-rich mergers/accretion scenario is generally consistent with the observed correlation between supermassive black holes (BH) and galaxy velocity dispersion or bulge mass \\citep{DiMatteo2005}. It is generally believed that the correlations indicate a joint evolution of galaxies and BHs, with BH growth happening at the same time as the bulge growth, and providing a self-regulation via feedback (\\citealt{Silk1998}; but see \\citealt{Peng2007} and \\citealt{Jahnke2011} for a non-causal origin of the correlations). We demonstrate that indeed the $\\sigma$ variation directly traces the bulge growth  \\citep{Kormendy2001}. In fact $\\sigma$ traces the central galaxy concentration and the bulge size as estimated on optical morphology, while the outer galaxy disks remain flat. This implies that, if the BH indeed accretes from the same gas that grows the bulge, BH mass should correlate better with $\\sigma$ than with total mass as observed \\citep{Gebhardt2000bh,Ferrarese2000}.  However, when ETGs experience dry mergers, their BH grows in proportion to the mass, but galaxies move along lines of nearly constant $\\sigma$. For this reason, an expectation from this picture is that BHs at the high mass end should start to appear too massive with respect to the predictions of the $M_{\\rm BH}-\\sigma$ relation \\citep[see also][]{Nipoti2003,Boylan-Kolchin2006,Ciotti2009}. The high-mass BH end is still not sufficiently populated to reliably test this prediction, but indirect evidence seems to support this possibility \\citep{Lauer2007}. Early direct evidences come from the recent detection of two giant black holes at the centre of two bright cluster galaxies, which are clearly more massive than the $M_{\\rm BH}-\\sigma$ prediction \\citep{McConnell2011} or the related observation that non-core ETGs follow a more shallow $M_{\\rm BH}-\\sigma$ relation than core ETGs \\citep{Graham2012}.  In summary we propose that the distribution of galaxy properties on the MP, where $M/L$ and age follows lines of constant $\\sigma$ on the MP, could be explained by the combination of two processes, which can happen multiple times during the evolution of a single galaxy: (i) They accrete gas, which grows the bulge and BH, shrinks their \\re, and increases $\\sigma_{\\rm e}$ and concentration, while some process which seems associated with the bulge growth (e.g.\\ AGN feedback), quenches star formation; (ii) They experience mostly minor dry mergers that move them along lines of roughly constant $\\sigma$ (or steeper).  An open question in the scenario in which ETGs evolve relatively quietly from spirals comes from the comparison of our findings with the empirical scaling relations one observes at larger redshift. In fact at $z\\ga1.5$ galaxies are found to be smaller than local ones with the same mass \\citep{Daddi2005,Trujillo2006,Trujillo2007,vanDokkum2008,Cimatti2008}. They populate the region below the ZOE of local galaxies, although their $\\sigma$ tend to be consistent with our local observations \\citep{Cappellari2012iau}. This may indicate that the compact high-redshift ETGs follow a different and more violent evolutionary path than the more quiet majority of local ETGs, as suggested by other high-redshift observations \\citep{Barro2012}. A way to reconcile this very mild (or lack of) $\\sigma$ evolution is by assuming that the compact primordial ETGs grow mostly by accretion of small satellites in their outer halos, while preserving the central structure \\citep{Naab2009,Hopkins2010,Oser2010,Oser2012}. This seems consistent with the shape of the photometric profiles of the early ETGs \\citep{Hopkins2009compact,Bezanson2009,vanDokkum2010profiles,Hilz2012b}. A caveat is that significant biases may still exist in the high-redshift photometry \\citep{Mancini2010}, considering that systematic differences of up to a factor of two exists even on well observed ETGs in the nearby Universe (\\citealt{Kormendy2009}, \\citealt{Chen2010}, Paper~I). Moreover comparisons of photometric profiles tend to be made against bona fide ellipticals, while the remnant of the high-redshift ETGs are likely disk-like fast rotators ETG, which have systematically different profiles. Finally the comparisons should ideally be done in mass density, instead of surface density, but kinematic information is available for only a handful of galaxies. This implies that there is perhaps still some room for the compact high-$z$ ETGs to become more consistent with local ones, than currently assumed."
    },
    "1208/1208.3009_arXiv.txt": {
        "abstract": "We investigate the variation with light quark mass of the mass of the nucleon as well as the masses of the mesons commonly used in a one-boson-exchange model of the nucleon-nucleon force. Care is taken to evaluate the meson mass shifts at the kinematic point relevant to that problem. Using these results, the corresponding changes in the energy of the $^1$S$_0$ anti-bound state, the binding energies of the deuteron, triton and selected finite nuclei are evaluated using a one-boson exchange model. The results are discussed in the context of possible corrections to the standard scenario for big bang nucleosynthesis in the case where, as suggested by recent observations of quasar absorption spectra, the quark masses may have changed over the age of the Universe. ",
        "introduction": "In the last decade there has been considerable interest in the possibility that the fundamental ``constants'' of Nature may actually change with time~\\cite{Uzan:2002vq}. Although it remains controversial, there is growing evidence that the fine structure constant may have varied by an amount of order a few parts in $10^{-5}$ over a period of 5--10 billion years~\\cite{Dzuba:1999zz,Webb:1998cq,Murphy:2003hw,Sandvik:2001rv}. It has even been suggested that this variation may have a dipole structure as we look back in different directions~\\cite{Webb:2010hc}. Although this possible variation is quite small, within the framework of most attempts at grand unification, a variation of $\\alpha$ implies considerably larger percentage changes in quantities such as $\\Lambda_{\\rm QCD}$ and in the quark masses~\\cite{SS,Bekenstein:1982eu,Olive:2002tz}. For example, in Ref.~\\cite{SS} it was shown that the variation $\\delta m_q/m_q$ would be of order 38 times that of $\\delta \\alpha/\\alpha$.  In the light of these developments it is very natural to ask what other signatures there may be for such changes. These may, for example, be the consequent changes in hadron masses or magnetic moments~\\cite{Flambaum:2002wq,Flambaum:2004tm,Cloet:2008wg}. Indeed, in some cases the level of precision possible in modern atomic, molecular and optical physics means that it may even be feasible to detect the minute variations expected under the hypothesis of linear variation until the present day over a period as short as a year~\\cite{Prestage:1995zz,Karshenboim:2000rg,Marion:2002iw}.  Another consequence of a variation in the parameters relevant to hadron structure is the possibility of observable consequences in big bang nucleosynthesis (BBN) or other nuclear phenomena such as the composition of the ash of long extinct natural nuclear reactors~\\cite{Shlyakhter_76,Damour:1996zw,FOth}. In this context, the effect of quark mass changes on the nucleon-nucleon force has been studied in effective field theory~\\cite{Beane:2002vs,Epelbaum:2002gb}, most recently including constraints from lattice QCD~\\cite{Beane:2002xf,Soto:2011tb}. {}For the moment these lattice studies are at too high a quark mass to provide an accurate constraint~\\cite{Soto:2011tb,Young:2002ib}. In an alternative approach based on a more traditional model of the nucleon-nucleon force, the latest work of Flambaum and Wiringa~\\cite{FW} on this topic involved the study of the variation of nuclear binding with quark mass using the Argonne potential and Schwinger-Dyson estimates of the variation of meson masses.  In this work we employ a one boson-exchange (OBE) model of the nuclear force to calculate the variation with changes in the quark mass of the binding energies of selected finite nuclei as well as the energy of the $^1$S$_0$ anti-bound state  and the binding energies of the deuteron and triton. Apart from its intrinsic interest, this approach complements the work of Ref.~\\cite{FW} and a comparison of the two provides one way to gauge the possible model dependence of the variations reported. The method used here involves a detailed study of the variation of the mass of each of the mesons usually employed in a one-boson-exchange (OBE) picture of the nucleon-nucleon (NN) force. Care is taken to estimate this shift at the relevant kinematic point, not just at the real, on-shell meson mass or its pole position. These changes are then introduced into the quark-meson coupling (QMC) model for some light nuclei and a typical OBE model for the two nucleon systems and a Faddeev calculation of the triton.  In section II we examine, in turn, each of the mesons $\\sigma_0, \\, \\sigma_1, \\, \\omega, \\, \\rho, \\, \\pi$ and $\\eta$. Section III presents results for finite nuclei, while the two nucleon system and triton are discussed in section IV. The final section is reserved for some concluding remarks. ",
        "conclusions": "\\label{conclusion} We have calculated the variation of the binding energy of the deuteron, triton  and the $^1$S$_0$ anti-bound pole position, as well as the binding energy per nucleon for a number of light nuclei, with respect to variations in the light (average of $u$ and $d$) quark mass. The results, expressed in terms of a parameter $K_A$, defined by \\begin{equation} \\frac{\\delta BE(A)}{BE(A)} = K_A \\frac{\\delta m_q}{m_q} \\, , \\label{eq:KA} \\end{equation} are summarised in Table-\\ref{coeff_summary}. In order to determine these coefficients, we first calculated the change with quark mass of the mesons used in a typical one-boson-exchange treatment of the nucleon-nucleon force. Those results were summarised in Table~\\ref{meson_summary}. {}For each nucleus we calculated the rate of change of the binding energy with respect to the mass of each meson and the mass of the nucleon itself. The values of $K_A$ were obtained by combining the latter with the results in Table~\\ref{meson_summary}. \\begin{table}[htb] \\begin{center} \\caption{Coefficients $K_{A}$ summarising the rate of variation of the binding energies and the $^1$S$_0$ anti-bound state pole with respect to quark mass - see Eq.~(\\ref{eq:KA}).} \\label{coeff_summary} \\vspace{0.3 cm} \\begin{tabular}{lc}\\hline Nucleus &~~~ $K_{A}$  \\\\ \\hline \\hline $D$             & -0.912  \\\\ \\hline $T$           &  -0.979 \\\\ \\hline $E_P$   & -2.839 \\\\ \\hline $^{7}Li$ &    -2.571 \\\\ \\hline $^{12}C$ &  -1.438  \\\\\\hline $^{16}O$          & -1.082  \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  {}For the deuteron our result, $K_d = -0.91$, is very close to that reported by Flambaum and Wiringa~\\cite{FW} using the AV14 potential, namely $-0.84$. Similarly for the triton, our value $K_t=-0.89$ is very close to their value of $-0.98$. The closeness of these results for two rather different treatments of the NN force lends considerable confidence in their reliability. However, for the position of the $^1S_0$ anti-bound state our calculation differs considerably from that of Ref.~\\cite{FW}, taking the opposite sign. This suggests that this quantity may be rather more model dependent than has been realized hitherto.  In the case of light nuclei, the binding energies reported here were calculated in the quark-meson coupling (QMC) model, a relativistic mean-field model that takes into account the self-consistent response of the internal structure of the nucleon to these mean fields. Through the self-consistency, the model yields many-body~\\cite{Guichon:2004xg} or equivalently density-dependent interactions~\\cite{Guichon:2006er}. Indeed, the density dependent Skyrme forces derived from QMC have proven remarkably realistic~\\cite{Dutra:2012mb}. The values of $K_A$ deduced in this way for $^7$Li, $^{12}$C and $^{16}$O are reported in Eqs.~(\\ref{ELi7})-~(\\ref{EO16}). It is interesting that the value obtained for $^7$Li, namely $K_{^7Li} = -2.57$, is significantly larger than that reported in Ref.~\\cite{FW}, namely $-1.03$ (AV14) and $-1.50$ (AV18+UIX). These authors did suggest that the uncertainty on the value of $K$ could be as large as a factor of two and our value is consistent at that level. Clearly, this degree of variation calls for more investigation to see whether the model dependence can be reduced.  Our study of these variations of binding energies with quark mass is, of course, motivated by the possible effects on big bang nucleosynthesis (BBN). Amongst the many challenges there, the sizeable discrepancy in the abundance of $^7$Li with the latest photon-to-baryon ratio (post WMAP) is of particular interest. Figure~\\ref{7Li_variation_with_mq} illustrates the $^7$Li abundance calculated using the BBN code of Kawano~\\cite{Kawano:1992ua}, if one allows {\\it only} the binding energy of the deuteron  and the energy of the virtual $^1S_0$ state to change with quark mass. The curves correspond to the values of $K_d$ and $K_P$ calculated here (solid line) as well as the values used by Berengut {\\it et al.}~\\cite{Berengut:2009js} (dashed line). The substantial difference in slope means that while a 3\\% shift in $\\delta m_q/m_q$ would suffice to reproduce the empirical abundance using the values of Berengut {\\it et al.}, with our values this would require a huge change in quark mass. This simple example illustrates the importance of a complete study of the BBN problem including all of the consequences of a shift of quark mass within the current approach, which we leave for future work. \\begin{figure}[ht] \\includegraphics[scale=0.7]{Li7CSSM.eps} \\caption{(Color online) Abundance of $^{7}$Li with respect to changes in the quark mass in $p(n,\\gamma)d$ calculated in the same way as \\cite{Berengut:2009js} (dashed-red line) and using our results for $K_{D}$ and $K_{E_P}$ (continuous-blue line).} \\label{7Li_variation_with_mq} \\end{figure} Finally, we note that while the variation of the light quark masses should be most important, it will also be necessary to take into account the effect of a corresponding change in the strange quark mass, especially now that the strange quark sigma commutator seems to be under control~\\cite{Young:2009ps}."
    },
    "1208/1208.3715_arXiv.txt": {
        "abstract": "{In this paper, we revisit the Cardassian model in which the radiation energy component is included. It is important for early epoch when the radiation cannot be neglected because the equation of state (EoS) of the effective dark energy becomes time variable. Therefore, it is not equivalent to the quintessence model with a constant EoS anymore. This situation was almost overlooked in the literature. By using the recent released Union2 $557$ of type Ia supernovae (SN Ia), the baryon acoustic oscillation (BAO) from Sloan Digital Sky Survey and the WiggleZ data points, the full information of cosmic microwave background (CMB) measurement given by the seven-year Wilkinson Microwave Anisotropy Probe observation, we constrain the Cardassian model via the Markov Chain Monte Carlo (MCMC) method. A tight constraint is obtained: $n= -0.0479_{-    0.0732-    0.148}^{+    0.0730+    0.142}$ in $1,2\\sigma$ regions. The deviation of Cardassian model from quintessence model is shown in CMB anisotropic power spectra at high $l$'s parts due to the evolution of EoS. But it is about the order of $0.1/\\%$ which cannot be discriminated by current data sets. The Cardassian model is consistent with current cosmic observational data sets.} ",
        "introduction": "Since the discovery of current accelerated expansion of our Universe \\cite{ref:Riess98,ref:Perlmuter99}, a flood of models have been designed to explain this late time accelerated expansion phase. For the reviews, please see \\cite{ref:DEReview1,ref:DEReview2,ref:DEReview3,ref:DEReview4,ref:DEReview5,ref:DEReview6,ref:DEReview7}. In the phenomenological perspective, the accelerated expansion of our Universe can be realized through modifying the form of Friedmann equation via the introduction of an extra exotic energy component, dubbed dark energy, which has negative pressure, or by some possible modifications of gravity theory, say $f(R)$ and brane models etc. As a result, the conventional Friedmann equation can be modified into the form of \\begin{equation} H^2=f(\\rho), \\end{equation} where $f(\\rho)$ is a function of energy density $\\rho$ which may include dark matter and extra energy components, and $H$ is the Hubble parameter. The Cardassian model firstly proposed by Freese and Lewis \\cite{ref:Cardassian}, where the Friedmann equation was modified into the form of \\begin{equation} H^2=\\frac{8\\pi G}{3}\\rho+A\\rho^n\\label{eq:FDE} \\end{equation} to explain the current accelerated expansion of our Universe. Here $\\rho$ can be composed of conventional matter, i.e, cold dark matter, baryons and radiation. For the origin of $\\rho^n$ term, one can find several explanations \\cite{ref:cardterm}. It can mimic the brane model which include a power law term due to the embedding of our universe into a five dimensional bulk. And, the late time accelerated expansion of our Universe is because of the leakage of the gravity force at the large scale in the brane world. The Cardassian model can reduce to $\\Lambda$CDM model when $n=0$ and to conventional CDM model with vanishing cosmological constant when $A=0$. The new term $\\rho^n$ will dominate the energy component at the late epoch to provide an accelerated expansion, then the values of $n$ should be $<2/3$.  This model has been confronted by cosmic observations extensively, when the energy density is composed of cold dark matter and baryon, for the recent result please see \\cite{ref:cardassiancon} and references therein, where $n=w+1=-0.039^{+0.135}_{-0.153}$ was obtained by using SN Ia, BAO, CMB shift parameters and observational Hubble data. In this case, the effective equation of state (EoS) of effective dark energy (the second term of Eq. (\\ref{eq:FDE})) is $w^{eff}_{de}=n-1$. So it corresponds to quintessence dark energy model where the matter and dark energy are included only. However, it is not always true. When the radiation energy component is added in the Cardassian model, i.e. $\\rho=\\rho_m+\\rho_r$, it is not equivalent to the quintessence model anymore due to the time variable effective EoS \\cite{ref:Sen} \\begin{equation} w^{eff}_{de}=(n-1)+\\frac{n}{3}\\frac{\\rho_{r0}a^{-4}}{\\rho_{m0}a^{-3}+\\rho_{r0}a^{-4}}\\label{eq:GEoS} \\end{equation} where $\\rho_0$ and $\\rho_{m0}=\\rho_{c0}+\\rho_{b0}$ are the present energy density for radiation and matter and we have normalized the scale factor to $a_{0}=1$. At the early epoch after the recombination and before the last scattering surface, the radiation and the matter energy components are important, where the effective EoS $w^{eff}_{de}$ is not a constant. Therefore it is important to take the general form of EoS (\\ref{eq:GEoS}) into account when one constrains Cardassian model from CMB observations. However, in the literature for example \\cite{ref:cardassiancon}, this important thing was {\\it not} considered at all, with exception of \\cite{ref:Sen} (see also \\cite{ref:spectrum}), where the locations of the peaks of the CMB anisotropic power spectrum were used as cosmic constraint. Though the time variable effective EoS was considered in \\cite{ref:Sen}, only the positions of peaks of CMB power spectra were used as cosmic constraints. In that paper \\cite{ref:Sen}, by fixing the values of $n_s$ and giving a number of values of $n$, the authors gave the ranges of $\\Omega_{m0}$ and $h$. In fact, it is not enough to give a tight constraint to model parameter space. In Ref. \\cite{ref:spectrum}, the authors investigated the CMB TT and matter power spectra. But, they did not give the model parameter space from the cosmic observations.  In the literature, the Cardassian model was constrained by the so-called CMB shift parameters, i.e. $R$, $l_a$ and $z_\\ast$ which are obtained based on a $\\Lambda$CDM model, from high redshifts. The problem that was almost overlooked is that the values of the CMB shift parameters depend on the cosmological models, here $\\Lambda$CDM model. The potential logic is that since $\\Lambda$CDM model is a concordance model then any  model which deviates slightly from $\\Lambda$CDM model is a competitive model too. However it is non-proper when one uses cosmic observational data points to constrain any other cosmological models. Strictly speaking, it is just a test of the possible viability of a model \\footnote{In the sense of small deviation from $\\Lambda$CDM model.} not a constraint to model. This is because that the potential {\\it circular problem} is committed and that is what we try to avoid in the cosmological constraint issue. In fact, the CMB shift parameters would be different for different cosmological models due to different physics process around the last scattering surface, for example early dark energy model \\cite{ref:EDE} where the contribution from dark energy may not be neglected due to nontrivial equation of state (EoS) of dark energy. In this situation, it is dangerous to use data points derived in $\\Lambda$CDM model due to much departure from $\\Lambda$CDM model. One the other hand, the full CMB data points contain more information than the shift parameters. For instance, at late epoch when the dark energy is dominated, the gravitational potential becomes evolution. It affects to the anisotropy power spectra of CMB at large scale (low $l$ parts) due to the so-called integrated Sachs-Wolfe (ISW) effect which is sensitive to the properties of dark energy. Apparently, the CMB shift parameters do not include this information. So, one can expect a tight constraint to cosmological models when the full information for CMB is included. So, in this paper, we should take the observational data points from CMB directly not the derived model dependent CMB shift parameters.  To fill out the gap and to avoid the so-called circular problem, in this paper we shall use the type Ia supernova (SN Ia), baryon acoustic oscillations (BAO) and full WMAP-7yr data points to constrain the Cardassian model.  This paper is structured as follows. In section \\ref{sec:method}, we give a brief review of Cardassian model and present the comic observational data sets and constraint methodalogy used in this paper. Section \\ref{sec:conclusion} is the conclusion. ",
        "conclusions": "\\label{sec:conclusion}  In summary, in this paper we have revisited the Cardassian model where the radiation energy component is included. In this situation, the EoS of effective dark energy for the Cardassian model is not a constant but time variable. Then the Cardassian model is not equivalent to quintessence model anymore. We performed a global fitting on the cosmological parameters in Cardassian model by using a completely consistent analysis where the full information from WMAP-7yr released data points were involved in a consistent way. We find out that the Cardassian model is consistent with current cosmic observational data sets. The constrained results are shown in Tab. \\ref{tab:results}. The results show that the model parameter $n$ in $1,2\\sigma$ regions are $n= -0.0479_{-    0.0732-    0.148}^{+    0.0730+    0.142}$ which are tighter than that, $n=-0.039^{+0.135}_{-0.153}$, obtained recently in \\cite{ref:cardassiancon}. In Ref. \\cite{ref:Sen}, the authors had used the positions of CMB peaks to constrain the Cardassian model. They found out that the values of $n$ depends on the values of $n_s$. They also pointed out that for $n_s=1$ the $n=0$ ($\\Lambda$CDM model) is not included in the allowed region. In this paper, we have used the full information from WMAP-7yr results not the positions of peaks of CMB or the derived CMB shift parameters in $\\Lambda$CDM model to constrain the model parameter space. The results show that $n=0$ ($\\Lambda$CDM) case is included in the allowed region. It means that currently available data points from SN, BAO, HST and CMB cannot distinguish a Cardassian model from a $\\Lambda$CDM model. We also show the effect of time variable EoS of effective dark energy to the CMB power spectra for the Cardassian model by comparison to that for the quintessence model. We find out the deviation is about the order of $\\mathcal{O}(10^{-3})$ which cannot be discriminated by current CMB data sets."
    },
    "1208/1208.0716_arXiv.txt": {
        "abstract": "Using a Newtonian model of the Solar System with all 8 planets, we perform extensive tests on various symplectic integrators of high orders, searching for the best  splitting scheme for long term studies in the Solar System. These comparisons are made in Jacobi and Heliocentric coordinates and the implementation of the algorithms is fully detailed for practical use. We conclude that high order integrators should be privileged, with a preference for the new $(10,6,4)$ method of \\citep{Bla12}. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.5769_arXiv.txt": {
        "abstract": "Observations from near the Eureka station on Ellesmere Island, in the Canadian High Arctic at 80$^{\\circ}$ North, benefit from 24-hour darkness combined with dark skies and long cloud-free periods during the winter.  Our first astronomical surveys conducted at the site are aimed at transiting exoplanets; compared to mid-latitude sites, the continuous darkness during the Arctic winter greatly improves the survey's detection efficiency for longer-period transiting planets. We detail the design, construction, and testing of the first two instruments: a robotic telescope, and a set of very wide-field imaging cameras. The 0.5m Dunlap Institute Arctic Telescope has a 0.8-square-degree field of view and is designed to search for potentially habitable exoplanets around low-mass stars. The very wide field cameras have several-hundred-square-degree fields of view pointed at Polaris, are designed to search for transiting planets around bright stars, and were tested at the site in February 2012. Finally, we present a conceptual design for the Compound Arctic Telescope Survey (CATS), a multiplexed transient and transit search system which can produce a 10,000-square-degree snapshot image every few minutes throughout the Arctic winter. ",
        "introduction": "The continuous wintertime darkness at polar sites can greatly increase the detection efficiency of time-domain astronomy programs. These properties have encouraged the development of Antarctic optical surveys of the Southern sky with imagers and small telescopes\\cite{Rico2010, Daban2010}, such as the Gattini cameras\\cite{Moore2006,Moore2008,Moore2010} at Dome C ($75^{\\circ}$S) and Dome A ($80^{\\circ}$S), and the Chinese Small Telescope Array (CSTAR\\cite{Wang2011}) which performed long-term photometry on 10,000 stars in a 23 deg$^2$ region centered on the South Celestial Pole.  In this paper we describe developments towards wide-field photometric monitoring programs in the Canadian High Arctic, which offers 24-hour access to the Northern sky in the winter months. At a latitude of $80^{\\circ}$N on Ellesmere Island, Canada's Eureka research base supports the nearby Polar Environment Atmospheric Research Laboratory (PEARL), a facility situated on a 600-m high ridge and designed primarily for atmospheric studies. This ``Ridge Lab'' has been demonstrated to have excellent weather in the winter months\\footnote{clear enough conditions for differential photometry $\\sim$85\\% of the time, and photometric conditions $\\sim$50\\% of the winter\\cite{Steinbring2012}}, good seeing\\cite{Hickson2010}, and dark skies\\footnote{papers in these proceedings}.  \\begin{sidewaysfigure}[tbp] \\centering \\resizebox{\\textwidth}{!} { \\includegraphics{block_diagram3.pdf} } \\caption{A block diagram of the autonomous wide-field survey facility at the Ridge Lab site. The AWCams have been deployed on the Ridge Lab roof and have demonstrated remote astronomical observatory operation via the Internet. The areas shaded in red are heated and/or insulated; blue areas are allowed to operate at ambient temperature. The \"warm box\" allows separation of the dome and control electronics from the Ridge Lab building by up to several hundred meters.\\label{fig:block_diagram}} \\end{sidewaysfigure}  Our Arctic science instruments are based on two tracks: very-wide-field, small aperture monitoring for both bright exoplanet transit and transient events, and a larger survey telescope designed to search for transiting exoplanets in the habitable zone of $\\approx$10,000 M-dwarfs.  Figure \\ref{fig:block_diagram} details the major astronomical components installed at or planned for operation at the Ridge Lab site in the near future. The Ridge Lab has proved to be a remarkably good platform for astronomy, providing a base of operations, power, and connectivity to the South via the broadband link of the Canadian Network for the Detection of Atmospheric Change (CANDAC). The wide-field cameras have already been deployed to the site, providing an initial demonstration of reliable operation. The 0.5m Dunlap Institute Arctic Telescope (DIAT) is currently undergoing on-sky testing in New Mexico, along with ongoing cold-testing of the hardware components. In Table \\ref{tab:camera_specs} we outline the properties of the systems described in this paper, including a conceptual design for a much larger survey which would build on experience with the current systems.  The challenging environment at the Ridge Lab site requires careful attention to the survivability of equipment deployed there. Instruments must be capable of surviving -50$^{\\circ}$C temperatures and storm events depositing snow and ice with wind speeds as high as 40m/s. In general, the snow-accumulation events require manual intervention (a technician with a shovel and broom) to clear instruments and buildings, but the instruments we describe here are otherwise designed to operate autonomously. We note some specific ruggedization strategies for each instrument below.  The paper is organized as follows: in Section \\ref{sec:wf_cameras} we evaluate the performance of two wide-field cameras which we deployed at the Ridge Lab site in 2012; in Section \\ref{sec:diat} we describe the design and testing of a wide-field, 0.5m telescope which we have ruggedized for Arctic operations; and in Section \\ref{sec:cats} we conclude by presenting a conceptual design of a much larger extremely-wide-field, high-cadence sky survey instrument.  \\begin{table} \\label{tab:camera_specs}  \\begin{small} \\centering \\begin{tabular}{llll} & {\\bf AWCam systems (deployed)} & {\\bf DIAT (under testing)} & {\\bf CATS (conceptual)} \\\\ \\hline Field of view & 25.4$^{\\circ}$ \\& 40.8$^{\\circ}$ & 0.93$^{\\circ}$ & 120$^{\\circ}$\\\\ Apertures     & 71 mm \\& 42 mm & 0.5m & 71 mm\\\\ Targets       & Exoplanet transits of bright stars & Exoplanet transits of cool stars & Very-high-cadence sky survey \\\\ Cadence       & 15 seconds & 20 minutes & 1-5 minutes \\\\ \\\\ \\end{tabular}  \\end{small} \\caption{The specifications of the three survey instruments described in this paper.} \\end{table} ",
        "conclusions": "The Canadian High Arctic offers an opportunity for 24-hour access to a dark Northern sky, with excellent weather conditions and potentially very good seeing. Astronomical science operations with wide-field telescopes began in February 2012, and they are planned to continue with full-winter campaigns and the addition of the DIAT 0.5m robotic telescope. The CATS conceptual design could provide continuous imaging coverage of much of the Northern sky; coupled to an upgraded Dunlap Institute Arctic Telescope equipped with a low-resolution spectrograph, the system would provide an integrated, rapid High Arctic sky survey and follow-up facility."
    },
    "1208/1208.2194_arXiv.txt": {
        "abstract": "{Using a deprojection technique, we study the X-ray properties of the galaxy cluster Abell 1835 observed with {\\it Chandra}, including temperature, abundance, electron density, gas mass fraction, and total mass. A comparison with the results without deprojection shows that the properties do not change much. When we compare the results with those of {\\it XMM-Newton}, the difference between the temperature profiles derived from {\\it Chandra} and {\\it XMM-Newton} data still exists, even if the point-spread function effect of {\\it XMM-Newton} is corrected. To investigate the reasons for the difference, we used the spectra to cross-calibrate the temperatures. They show that the {\\it Chandra} spectra can be fitted well with {\\it XMM-Newton} temperatures. Furthermore, we derive the electron density profile from {\\it Chandra} data with {\\it XMM-Newton} temperatures and calculate the projected mass, which is consistent with the {\\it XMM-Newton} mass and a little lower than the weak lensing mass at $r_{500}$. Thus, it seems that the temperature derived from {\\it XMM-Newton} may be more reliable. \\\\ ",
        "introduction": "Galaxy clusters, the largest objects in the universe, originate in the primordial density perturbations from cosmic gravitational collapse. They are used in a variety of ways to perform both cosmological and astrophysical studies. Modern astronomy satellites such as {\\it XMM-Newton} and {\\it Chandra} have high sensitivity and spatial resolution, and unprecedented results have been gained through detailed analysis of their data. However, there are discrepancies between the properties of galaxy clusters derived from {\\it Chandra} and {\\it XMM-Newton} data, such as gas temperature and total mass. Nevalainen et al. (2010) examined the cross-calibration of the energy dependence and normalization of the effective area of {\\it Chandra} and {\\it XMM-Newton}, finding that the discrepancies of the $0.5-7.0$ keV band temperature measurements of galaxy clusters with {\\it EPIC/XMM-Newton} and {\\it ACIS/Chandra} could reach $\\sim10-15\\%$ on average.%  Abell 1835 is a classical bright cluster with a big cool core  (Allen et al. 1996). Its X-ray morphology (Schmidt et al. 2001; Smith et al. 2005) shows that it is an undisturbed and relaxed cluster. It is also an optimal candidate for a triaxial joint analysis via X-ray, SZ, and lensing techniques (Morandi et al. 2011). The temperature difference between the {\\it Chandra} and {\\it XMM-Newton} analysis also exists in Abell 1835: $\\sim12$ keV from {\\it Chandra} data (Schmidt et al. 2001) and $\\sim7.6$ keV from {\\it XMM-Newton} data (Jia et al. 2004). After the point-spread function (PSF) correction of {\\it XMM-Newton} data, Wang et al. (2010) derived a temperature profile similar to that of Jia et al. (2004). Therefore, the difference is not due to the PSF effect.  In addition, the masses of Abell 1835 derived from {\\it XMM-Newton} and {\\it Chandra} are different (Jia et al. 2004; Schmidt et al. 2002). Fortunately, the masses have been measured with the strong lensing method (Richard et al. 2010) as well as with the weak lensing method (Zhang et al. 2008). Since gravitational lensing directly probes the cluster total mass without any strong assumptions about the equilibrium state of the cluster, the lensing mass is generally more reliable.    Recently, the comparison of X-ray and gravitational lensing masses has been studied in detail by both observational and simulated analyses. Generally, the X-ray mass is consistent with or lower than the gravitational lensing mass. Richard et al. (2010) showed that the ratio of strong lensing mass and X-ray projected mass within $r<250$ kpc, $M_{SL}/M_X$, was 1.3. It was also found that the mass derived from the X-ray measurement is about half of the strong lensing mass in some clusters (e.g., A1689, Andersson \\& Madejski 2004, Lemze et al. 2008; PKS 0745-191, Chen et al. 2003). Moreover, Zhang et al. (2008) showed that the average ratio of the weak lensing mass to X-ray mass was $1.09\\pm0.08$, while Mahdavi et al. (2008) demonstrated that $M_X/M_{WL}$ is $1.03\\pm0.07$ and $0.78\\pm0.09$ at $r_{2500}$ and $r_{500}$, respectively. N-body/hydrodynamical simulation work also estimated the ratio between X-ray and lensing masses, $M_X/M_{WL}$,  which was $0.88\\pm0.02$ and $0.75\\pm0.02$ at $r_{500}$ in Meneghetti et al. (2010) and Rasia et al. (2012), respectively.  In this work, we reanalyse the {\\it Chandra} data of Abell 1835 with the same deprojection technique as Jia et al. (2004) to ascertain if the differences in temperatures derived from {\\it Chandra} and {\\it XMM-Newton} data can be corrected by data analysis. Furthermore, we investigate the reasons for the temperature differences and establish which result is more reasonable.  This paper is organized as follows: The {\\it Chandra} observations and data preparation are described in Section 2. Section 3 shows the basic spectra analysis, while Section 4 presents the deprojected electron density profile and the mass profile. In Section 5, we discuss the reasons for the temperature differences derived from {\\it Chandra} and {\\it XMM-Newton} data. We present our conclusions in Section 6.  Throughout this paper, we assume $H_0$ = 70 km s$^{-1}${Mpc}$^{-1}$, $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$. One arcminute corresponds to 236.2 kpc at Abell 1835 redshift of 0.2523. The selected energy band is 0.5 keV- 7.0 keV. ",
        "conclusions": "We have presented a detailed analysis of about 117 ksec of {\\it Chandra} observations on the galaxy cluster Abell 1835. Through the deprojected spectra analysis, we derived the deprojected temperatures, which do not differ much from the temperatures without deprojection and are still much higher than those of {\\it XMM-Newton}.   The total mass within the radius of $6'$, $M_{tot-Chandra}$ = 1.36$\\pm$0.42$\\times$10$^{15}$ M$_{\\odot}$, is much higher than the total mass of $M_{tot-XMM}$ = 0.84$\\pm$0.10$\\times$10$^{15}$ M$_{\\odot}$ derived from {\\it XMM-Newton} data. The difference of total mass is due to the discrepancy of the temperatures derived from {\\it Chandra} and {\\it XMM-Newton}. We also calculated the projected mass within the optical lensing arc, which is lower than the strong lensing mass derived by Richard et al. (2010).   After deprojection and PSF correction (Wang et al. 2010), the temperature and total mass of Abell 1835 resulting from {\\it Chandra} are still different from those of  {\\it XMM-Newton}. These differences may result from the calibration of the two instruments. We find that {\\it XMM-Newton} has a stronger restriction on temperature. And with {\\it XMM-Newton} temperatures, the projected mass from {\\it Chandra} data is lower than the weak lensing mass and consistent with other observational analyses. For these reasons, the temperatures obtained from {\\it XMM-Newton} may be more reliable."
    },
    "1208/1208.0158_arXiv.txt": {
        "abstract": "The discovery of a low-mass star with such low metallicity as $\\le 4.5 \\times 10^{-5}$ $Z_\\odot$ reveals the critical role of dust in the formation of extremely metal-poor stars. In this paper we explore the effect of the growth of dust grains through accretion of gaseous refractory elements in very low-metallicity pre-stellar cores on the cloud fragmentation induced by the dust emission cooling. Employing a simple model of grain growth in a gravitationally collapsing gas, we show that Fe and Si grains can grow efficiently at hydrogen densities of $\\simeq 10^{10}$--$10^{14}$ cm$^{-3}$ in the clouds with metal abundances of $-5 \\la$ [Fe, Si/H] $\\la -3$. The critical metal number abundances, above which the grain growth could induce the fragmentation of the gas clouds, are estimated to be $A_{\\rm crit} \\simeq$ $10^{-9}$--$10^{-8}$, unless the initial grain radius is too large ($\\ga$ 1 $\\mu$m) or the sticking probability is too small ($\\la$ 0.01). We find that even if the initial dust-to-gas mass ratio is well below the minimum value required for the dust-induced fragmentation, the grain growth increases the dust mass high enough to cause the gas fragmentation into sub-solar mass clumps. We suggest that as long as the critical metal abundance is satisfied, the grain growth could play an important role in the formation of low-mass stars with metallicity as low as $10^{-5}$ $Z_\\odot$. ",
        "introduction": "Dust grains in the early universe are considered to be important agents to trigger the formation of low-mass stars in metal-poor environments (Omukai 2000; Schneider et al.\\ 2003); the cooling of the gas through thermal emission of dust makes collapsing dense cores gravitationally unstable, leading to the fragmentation into multiple sub-solar mass clumps at gas densities of $10^{12}$--$10^{14}$ cm$^{-3}$ (Tsuribe \\& Omukai 2006; Dopcke et al.\\ 2011). This scenario has been recently supported by the discovery of a Galactic low-mass star, SDSS J102915+172927 (Caffau et al.\\ 2011), whose metal content is too low ($Z \\le 4.5 \\times 10^{-5}$ $Z_\\odot$) to induce the fragmentation of star-forming clouds by metal-line cooling (see Klessen et al.\\ 2012; Schneider et al.\\ 2012b for details).  The condition that realizes the dust-induced fragmentation depends on the amount of dust grains as well as their size distribution in pre-stellar clouds (Omukai et al.\\ 2005; Schneider et al.\\ 2006, 2012a). Schneider et al.\\ (2012a) found that the formation condition of the low-mass fragments obtained by the numerical simulations is fully described in terms of the product of dust-to-gas mass ratio $\\mathcal{D}$ and geometrical cross section per unit dust mass $\\mathcal{S}$ as follows; \\begin{eqnarray} \\mathcal{S D} > 1.4 \\times 10^{-3} ~ {\\rm cm}^2 ~ {\\rm g}^{-1} \\left( \\frac{T_{\\rm gas}}{10^3 ~ {\\rm K}} \\right)^{-\\frac{1}{2}} \\left( \\frac{c_{\\rm H}}{10^{12} ~ {\\rm cm}^{-3}} \\right)^{-\\frac{1}{2}}, \\end{eqnarray} where $T_{\\rm gas}$ is the temperature of the gas, and $c_{\\rm H}$ is the hydrogen number density. Treating self-consistently the dust formation in the ejecta of supernovae (SNe) and the subsequent destruction of the dust by the reverse shocks, Schneider et al.\\ (2012a) argued that the condition could not be satisfied in the collapsing star-forming clouds enriched with metals and dust from the first SNe if a majority of grain formed in the SN ejecta are destroyed by the reverse shock (see also Schneider et al.\\ 2012b). However, it could be possible that the accretion of gaseous refractory elements released from dust grains in the shocked gas onto the surfaces of the SN dust surviving in star-forming clouds changes the mass and size distribution of the dust, and thus affects the thermal evolution of the collapsing cores.\\footnote{Hirashita \\& Omukai (2009) examined the coagulation of dust in collapsing clouds with a variety of metallicity. They found that the dust coagulation can proceed even at metallicity as low as $10^{-6}$ $Z_\\odot$ but does not have any impact on the thermal evolution of the star-forming clouds.}  In this paper, we investigate the feasibility of grain growth in low-metallicity star-forming clouds to explore whether the grain growth can facilitate the formation of metal-poor low-mass stars. In Section 2, we describe the model of grain growth in collapsing dense clouds, and present the results of the calculations in Section 3. In Section 4, we estimate the critical metal abundances above which the grain growth could encourage the gas fragmentation into sub-solar mass clumps, and discuss the corresponding dust-to-gas mass ratio and total metallicity. The conclusion is given in Section 5. ",
        "conclusions": "We have investigated the growth of dust grains in metal-poor proto-stellar clouds. Our simple model shows that the grain growth can operate efficiently even in collapsing dense cores with metal abundances as low as [Fe, Si/H] $\\simeq -5$. We also present the critical metal abundances above which the grain growth could affect the fragmentation process of collapsing gas clouds. This abundance is estimated to be $A_{\\rm crit} \\simeq$ $10^{-9}$--$10^{-8}$, which suggests that the formation of low-mass stars with metallicity of $\\sim$$10^{-5}$ $Z_\\odot$ can be possible. We conclude that even if the initial dust-to-gas mass ratio does not satisfy the condition required for the dust-induced fragmentation, the grain growth can increase the dust-to-gas mass ratio high enough to facilitate the formation of metal-poor low-mass stars.  Our results suggest that if grain growth is considered, the formation of low-mass protostars can occur not only at very low metallicity but also at higher metallicity. The final mass of a newly born star is determined by the accretion of the surrounding gas onto the protostars (McKee \\& Ostriker 2007 and references therein), and its mass accretion rate would be regulated by the mass of the central protostar induced by the grain growth. Thus, the grain growth in collapsing clouds might be a fundamental physical process to control the stellar initial mass function in the present universe.  It should be mentioned that we have considered only the growth of single-component Fe and Si grains with a single initial radius. However, it might be possible that Si and Fe atoms condense as silicates or oxides in an oxygen-rich gas. Since the growth of such compound grains with no monomer molecule has been usually treated by considering Si or Fe element as a key element (e.g., Zhukovska et al.\\ 2008), the mass and radius of dust given in this paper are considered to be lower limits. On the other hand, the timescale of grain growth is sensitive to the initial grain radius (Hirashita \\& Kuo 2011). Thus, the effect of the initial size distribution as well as the growth of compound grains should be explored. Furthermore, we have assumed the sticking probability of $s_i = 1$ and a constant gas temperature $T_{\\rm gas} = 10^3$ K during the collapse of the clouds. In particular, too low sticking probabilities ($s_i \\la 0.01$) may prevent the grain growth from becoming efficient for the metal abundances of $A_i \\la 2.5 \\times 10^{-7}$. We note that our conclusions obtained by the simple model should be confirmed by more sophisticated simulations of the thermal evolution of star-forming clouds involving grain growth."
    },
    "1208/1208.2688_arXiv.txt": {
        "abstract": "\\noindent We observed SN~2002bu in the near-IR with the Hubble Space Telescope, the mid-IR with the Spitzer Space Telescope and in X-rays with Swift 10 years after the explosion.  If the faint $L_H \\sim 10^2 L_\\odot$ HST near-IR source at the transient position is the near-IR counterpart of SN~2002bu, then the source has dramatically faded between 2004 and 2012, from $L \\simeq 10^{6.0}L_\\odot$ to $L \\simeq 10^{4.5}L_\\odot$. It is still heavily obscured, $\\tau_V \\simeq 5$ in graphitic dust models, with almost all the energy radiated in the mid-IR. The radius of the dust emission is increasing as $R \\propto t^{0.7 \\pm 0.4}$ and the optical depth is dropping as $\\tau_V \\propto t^{-1.3\\pm0.4}$. The evolution expected for an expanding shell of material, $\\tau_V \\propto t^{-2}$, is ruled out at approximately $2\\sigma$ while the $\\tau_V \\propto t^{-0.8}$ to $t^{-1}$ optical depth scaling for a shock passing through a pre-existing wind is consistent with the data.  If the near-IR source is a chance superposition, the present day source can be moderately more luminous, significantly more obscured and evolving more slowly.  While we failed to detect X-ray emission, the X-ray flux limits are consistent with the present day emissions being powered by an expanding shock wave. SN~2002bu is clearly a member of the SN~2008S class of transients, but continued monitoring of the evolution of the spectral energy distribution is needed to conclusively determine the nature of the transient. ",
        "introduction": "\\label{sec:introduction}  Supernova (SN) 2002bu was discovered on 2002 March 28 by \\cite{Puckett2002} in the galaxy NGC~4242, with a fairly low peak magnitude $M_V \\simeq -15$ \\citep{Hornoch2002}. The last pre-discovery observations on 2001 February 21 and March 14 placed a unfiltered magnitude limit at the location of the  SN of 20.5~mag. A low resolution spectrum taken by \\cite{Ayani2002} on April 1 showed a flat continuum and strong, narrow Balmer emission lines (FWHM $\\simeq$ 1100 km/s) that led to a Type~IIn classification of the SN. The early ($\\sim$80 days post peak) light curve presented by \\cite{Foley2007} and \\cite{Smith2011} showed a plateau and a slower rise and decline as compared to other Type~II SNe, with the color becoming redder with time.  \\cite{Thompson2009} proposed that SN~2002bu may be a member of a new class of stellar transients, with the prototypes being SN~2008S \\citep{Arbour2008} and the 2008 optical transient (OT) in NGC~300 \\citep{Monard2008}. The class is characterized by a transient--progenitor pair, where the transient is $2-3$ mag fainter than a regular core collapse SN (ccSN), with narrow emission lines and evidence of internal extinction in the spectrum, while the progenitor is a dust enshrouded, low luminosity ($\\sim$5$\\times 10^4 L_{\\odot}$) star with little variability in the years before the outburst (\\citealt{Prieto2008a}, \\citealt{Prieto2008b}). In mid--IR color--magnitude diagrams (CMD), the progenitors occupy the extreme end of the asymptotic giant branch (AGB) sequence and the obscuring dust is graphitic rather than silicate (\\citealt{Prieto2009}, \\citealt{Wesson2010}). \\cite{Thompson2009} further showed that these transients are relatively common among ccSNe ($\\sim$20\\%), but their progenitors are extremely rare ($\\ltorder 10^{-4}$ of massive evolved stars), which implies that many massive stars go through this dust obscured phase shortly ($\\leq10^{4}$ yrs) before the explosion.  The rarity of the progenitor stars was further confirmed in the survey of additional galaxies by \\cite{Khan2010}.  Since the main characteristic of the new class is a dust enshrouded progenitor, it is impossible to unambiguously classify SN~2002bu as a member, as there are no pre--explosion IR observations of the region. However, \\cite{Thompson2009} found a bright, red mid-IR source at the location of SN~2002bu in Spitzer Space Telescope (SST) data taken 2 years after the explosion, which indicates dust formation. This, together with the transient characteristics (low luminosity, dust and narrow emission lines visible in the spectrum), make SN~2002bu a likely member of the new class. In addition, \\cite{Smith2011} analyzed 5 spectra taken between 11 and 81 days after discovery, and they show that the spectrum of SN~2002bu becomes redder with time and evolves from a spectrum resembling an LBV in outburst at early times, to one more similar to the spectra of SN~2008S and the NGC~300~OT 81 days later. The evolution of the H$\\alpha$ line from a Lorentzian profile to an asymmetric, blueshifted Gaussian profile, is suggestive of newly formed dust. In \\cite{Kochanek2012} we argued that the early light curve also indicates dust formation, albeit with some ambiguities.  In \\cite{Kochanek2012} we analyzed archival Hubble Space Telescope (HST) and SST data from approximately two years after the transient peak.  As previously reported by \\cite{Thompson2009}, SN~2002bu was a luminous mid-IR source in 2004, but \\cite{Kochanek2012} also found that it was invisible in archival HST data from 2005 in the BVRI bands to limits of $\\sim 25$~mag.  The spectral energy distribution (SED) in 2004 was well fit by surrounding a $T_{*}=20000$~K, $L_{*}\\simeq10^{5.9} L_{\\odot}$ source with an optically thick ($\\tau_{V}\\simeq30$) shell of dust located at a radius of approximately $R\\simeq 10^{15.8}$~cm.  \\begin{figure} \\plotone{fig1.ps} \\caption{The top two panels show the [3.6] $\\mu$m SST observations of SN~2002bu from April 2004 (left) and January 2012 (right), the middle two panels show the corresponding [4.5] $\\mu$m SST data, and the bottom two panels show the  F110W ($J$) and F160W ($H$) HST   observations from February 2012 of the region marked with a square on the SST images. The four top panels are 12\\farcs0 $\\times$ 12\\farcs0. The two bottom panels are 3\\farcs0 $\\times$ 3\\farcs0 (about 80 $\\times$ 80 parsecs) and the radius of the circle is $\\sim$0\\farcs3, which is 3 times the uncertainty in the astrometry. } \\label{fig:ssthst} \\end{figure}  Broadly speaking, there are three possible explanations of the observations, as we discussed in our mid-IR survey of the ``supernova impostors''(\\citealt{Kochanek2012}). The first, ``traditional'', view of these events is that a shell of material is ejected during the optical transient and forms dust once the ejected material becomes cool enough.  As the shell expands, its optical depth drops as $\\tau \\propto 1/t^2$ and the characteristic temperature drops as $T \\propto L(t)^{1/4} t^{-1/2}$ because the shell expands with the ejecta velocity $R \\simeq v_s t$.  The dust radius at the time of the previous SST observations of SN~2002bu was consistent with a shell expanding at the velocity of $v_s=893$~km/s adopted by \\cite{Smith2011}. The second possibility is that the optical transient is a signal that the star is entering a high mass loss phase with a dense wind that forms dust and obscures the source.  While the wind is steady, the optical depth is roughly constant with a dust temperature close to the dust destruction temperature $T_d \\simeq 1500$~K unless the wind becomes optically thick in the mid-IR. When the high mass loss phase ends, the evolution quickly resembles the first scenario (see \\citealt{Kochanek2012}). The third scenario is the one introduced by \\cite{Kochanek2011} to explain SN~2008S and the NGC~300~OT.   Here the progenitors are already shrouded by a very dense wind when an explosive transient occurs that destroys most of the dust to leave the transient little obscured at peak.  The wind is so dense, however, that the dust reforms and re-obscures the transient.  The present day luminosity is a combination of a surviving star (if any, nothing in the data requires one) and the luminosity generated by the shock propagating through the wind.  At later times, the optical depth outside the shock is dropping as $\\tau \\propto 1/t$ and, once the optical depth is low enough, the X-rays produced in the shock should be observable.  These scenarios make different predictions for the time evolution of the transient, so monitoring the evolution of SN~2002bu as a function of wavelength should reveal the nature of the event.  Here we report near-IR HST, mid-IR SST and Swift X-ray observations of SN~2002bu taken roughly 10 years after the transient peak and 8 years after the the last HST and SST observations. Section~2 presents the new observational data.  We discuss the results in Section~3, and consider their broader implications in Section~4. ",
        "conclusions": "\\begin{figure*} \\plotone{fig4.ps} \\caption{The top left panel shows the SN~2002bu SED roughly 2 years after the explosion, while the other 4 panels show the current SED for sources A, B, and C (see Fig.~\\ref{fig:closeup}), in the case of source A also treating HST near-IR measurements as upper limits (top right panel).  Magnitudes from Table~\\ref{tab:mags} are converted to fluxes, and then to luminosities as $L=4\\pi D^2 \\nu  F_\\nu$ where $D=5.8$~Mpc. The SED models are black bodies with $A_V=0.012$~mag of total (Galactic) extinction. The thick black curves show the probability averaged SEDs for graphitic (solid) and silicate (dashed) dust, while the thin curves show the spread around the mean values. In the case of the 2004 SED, solid and dashed lines are the best fit graphitic and silicate models taken from \\cite{Kochanek2012} and correspond to a $T_{*}=20000$~K, $L_{*}\\simeq10^{5.9} L_{\\odot}$ star surrounded by an optically thick ($\\tau_{V}\\simeq30$) shell of dust expanding at $\\sim$900~km/s. } \\label{fig:sed} \\end{figure*}   We used DUSTY (\\citealt{Ivezic1997}; \\citealt{Ivezic1999}) to fit the SEDs of the sources detected in HST near-IR data. Figure~\\ref{fig:sed} shows the probability averaged SEDs of SN~2002bu in May~2004 and January~2012 (for sources A, B and C, treating near-IR HST measurements of source A either as detections or upper limits). The best fit model to the 2004/2005 data is taken from \\cite{Kochanek2012} and corresponds to a $T_{*}=20000$~K, $L_{*}\\simeq10^{5.9} L_{\\odot}$ star surrounded by an optically thick ($\\tau_{V}\\simeq30$) shell of dust expanding at $\\sim$900~km/s. Both graphitic and silicate models give similarly good fits to the data.  We also embedded DUSTY in a Markov Chain Monte Carlo engine to model the SEDs, varying the dust temperature, $T_d$, optical depth, $\\tau_V$, and stellar temperature, $T_*$, with either a fixed 2:1 ratio between the inner and outer radii of the shell or allowing the ratio to vary between 1.1 and 10.  We included a weak prior on the stellar temperature, $\\log_{10} T_* = 4.0 \\pm 0.3$, (and restricted its range to $3000~\\hbox{K} < T_* < 30000~\\hbox{K}$) and on the expansion velocity implied by the inner radius of the shell, $\\log_{10}(v_s/\\hbox{km/s})=\\log_{10}(893)\\pm0.30=2.95\\pm0.30$.  For the 2012 epoch we ran models considering star A as either a detection or an upper limit.  The results allowing for variations in the shell thickness were little different from those with a fixed thickness and no particular thickness was preferred, so we only report the results for the fixed thickness in Table~\\ref{tab:results}.  The results for the 2004 epoch are the same as in our earlier models from \\cite{Kochanek2012}.  The source must be quite luminous, $L_* \\simeq 10^{5.92 \\pm 0.02}L_\\odot$, but with an indeterminate source temperature because of the heavy obscuration.  The optical depth in graphitic models is $\\log_{10} \\tau_V \\simeq 1.6 \\pm 0.2$ (scattering plus absorption), and it is moderately higher in the silicate models because of the higher scattering opacities of silicate dusts (see the discussion in \\cite{Kochanek2012}). Significantly higher optical depths begin to have significant opacity even for the shorter wavelength IRAC bands, inconsistent with the shape of the SED. The dust temperatures at the inner edge of $T_d \\simeq 1100 \\pm 170$~K are roughly in the range expected for newly forming dust. The dust radius is estimated to be $\\log_{10} (R_{in}/\\hbox{cm}) \\simeq 15.77 \\pm 0.12$, implying a velocity of $\\log_{10} (v_s/\\hbox{km/s}) \\simeq 2.95 \\pm 0.12$ that is consistent with the velocity prior taken from the line widths cited by \\cite{Smith2011}. Note, however, the uncertainties are much smaller than those of the prior, so the data are in fact determining this radius and velocity rather than the prior.  Despite the reasonable coverage of the mid-IR SED, graphitic and silicate dusts fit the data equally well.  As discussed earlier, the primary effect of distance uncertainties is simply to rescale the luminosities, velocities and distances.  In particular, adopting the distance to NGC~4258 ($7.2$~Mpc) raises the velocity from $v_s \\simeq 900$~km/s to $v_s \\simeq 1100$~km/s, while using the \\cite{Springob2009} distance of $10.4$~Mpc raises it to $v_s \\simeq 1500$~km/s.  One reason we adopted the smaller distance is that the larger distances begin to require significantly larger mean expansion velocities than implied by spectroscopic observations during the transient (see \\citealt{Smith2011}).  The interpretation of the SED in 2012 depends critically on whether any source in the HST images corresponds to SN~2002bu.  If source A (or B/C) is a detection, then the present day source must be a relatively hot, $T_* > 15000$~K, star.  For graphitic dust, the star is relatively low luminosity $L_*=10^{4.5\\pm0.2} L_\\odot$, with moderate $\\log_{10} \\tau_V \\simeq 0.70\\pm0.25$ obscuration, at a relatively large radius $\\log_{10} (R_{in}/\\hbox{cm}) \\simeq 16.18\\pm0.25$ that corresponds to a lower average expansion velocity $\\log_{10} (v_s/\\hbox{km/s}) \\simeq 2.72\\pm 0.25$ but is uncertain enough to be consistent with no change in velocity. The dust temperature at the inner edge is much cooler, $T_d \\simeq 500$~K. In the silicate models the source is more luminous, $L_*=10^{4.95\\pm0.18}L_\\odot$ with greater optical depth, $\\log_{10} \\tau_V \\simeq 1.21\\pm0.12$, located at a smaller radius $\\log_{10} (R_{in}/\\hbox{cm}) \\simeq 15.98\\pm0.17$, that now requires a slower expansion rate $\\log_{10} (v_s/\\hbox{km/s})\\simeq 2.52 \\pm 0.17$.  The dust  temperature at the inner edge is moderately warmer, $T_d \\simeq 600$~K. There is no basis for choosing between the two dust types. If the association with source A is simply a coincidence and we instead treat the HST fluxes as upper limits, then solutions with cool stars are allowed, but they require increasingly high optical depths for lower temperatures, $\\tau_V \\simeq 100$ for the coolest ($T_*=2500$~K) models. This is essentially the limit of the fully obscured progenitors of SN~2008S and the NGC~300-OT (\\citealt{Prieto2008a}, \\citealt{Prieto2008b}, \\citealt{Thompson2009}, \\citealt{Kochanek2011}).   \\begin{figure} \\plotone{fig5.ps} \\caption{The mid-IR CMD of NGC~2403 (\\citealt{Khan2010}). The color evolution of SN~2002bu is marked with black squares connected by a dotted line, where the top point corresponds to the 2004 epoch, while the next to 2011 and 2012 epochs (see Table~\\ref{tab:mags}). For comparison, we show the color evolution of SN~2008S from 2008 to 2012 (blue circles)  and NGC~300-OT from 2008 to 2012 (red triangles) as summarized in \\cite{Kochanek2011} and \\cite{Szczygiel2012}. } \\label{fig:midIR_cmd} \\end{figure}  The current position of the source in the mid-IR CMD (see Figure~\\ref{fig:midIR_cmd}), close to the tip of the AGB sequence, is similar to that of the SN~2008S or NGC~300-OT progenitor stars (\\citealt{Prieto2008a}, \\citealt{Prieto2008b}).  This is simply a coincidence driven by the present luminosity and dust radius, because in steady state a surviving star could not support the necessary optical depths given the velocities needed to have new material at such large distances. The population studies by \\cite{Thompson2009} and \\cite{Khan2010} imply a lifetime in the obscured phase of order $t_o \\simeq 10^4$~years.  The mass loss rate is related to the optical depth by \\begin{equation} \\tau_V = { \\dot{M} \\kappa_V \\over 4 \\pi v_w R_{in} }, \\label{eqn:opdepth} \\end{equation} for visual opacity $\\kappa_V = 100 \\kappa_2$~cm$^2$/g, mass loss rate $\\dot{M}$ and wind speed $v_w$, which implies a total mass loss of \\begin{eqnarray} \\dot{M} t_o &\\simeq  &4 \\pi v_w R_{in} \\tau_v t_o \\kappa_V^{-1} \\\\ &\\simeq &{200\\over \\kappa_2} \\left( { v_w \\over 10^3~\\hbox{km/s} }\\right) \\left( { R_{in} \\over 10^{16}~\\hbox{cm} } \\right)~ \\left( { \\tau_V \\over 10 } \\right) \\left( { t_o \\over 10^{4}~\\hbox{years} }\\right) M_\\odot \\nonumber \\end{eqnarray} while the star is in its obscured phase.  No star could sustain this for $v_w \\sim 10^3$~km/s, so the current state cannot represent a long lived period and must be a transient phase.  The progenitors, under the super-AGB star hypothesis of \\cite{Thompson2009}, can have the necessary lifetimes because dust driven wind velocities of $v_w \\simeq 10$-$20$~km/s are so low (e.g. \\citealt{Ivezic2010}).  However, such a slow wind commencing after the transient would not even have started to form dust at this point in time.   \\begin{figure*} \\plotone{fig6.ps} \\caption{The average expansion velocity and optical depth of the DUSTY models for SN~2002bu for graphitic (left) and silicate (right) dusts.  The black (red) contours encompass the $1\\sigma$ region for two parameters for the 2004 (2012) epoch.  The dashed green contours (labeled shock) show the expected parameters in 2012 given the parameters in 2004 assuming the radius is expanding as $R\\propto t^x$ with $x=0.8$ (left) or $1.0$ (right) and the optical depth is dropping as $\\tau_V \\propto 1/R$.  The dotted blue (labeled shell) shows the expected parameters in 2012 if the radius is expanding as $R \\propto t$ and the optical depth is dropping as $\\tau_V \\propto 1/t^2$.  Shock models to the right of the heavy solid line are excluded ($1\\sigma$) by the upper limits on the X-ray flux in the Swift observations.  If the HST sources are treated as upper limits, the 2012 contours expand upwards to higher optical depths.  The optical depth is related to the wind density by Eqn.~\\ref{eqn:opdepth} where we used $\\kappa_2=1.7$ ($0.84$) for graphitic (silicate) dusts. } \\label{fig:params} \\end{figure*}  The second possibility is the traditional view of the supernova impostors as stars that briefly enter a high mass loss state during the transient to produce a shell of ejected material that then forms dust.  As discussed in \\cite{Kochanek2012}, SN~2002bu probably started forming dust too soon after the transient peak to have formed it in ejected material.  The more serious problem is that between 2004 and 2012 an expanding, dense shell should have a radius growing as $R_{in} \\propto t$ and an optical depth dropping as $\\tau_V \\propto t^{-2}$.  If we take the elapsed time from discovery to the IRAC observations in 2004 (736 days) and 2012 (3566 days), the radius should have increased by a factor of $4.9$ and the optical depth should have dropped by a factor of $23.5$.  Fig.~\\ref{fig:params} illustrates this by comparing the mean expansion velocities and optical depths for the two epochs and both grain types.  If we scale the values in 2004 to their expected values in 2012 simply using the expected temporal scalings, we see that the observed properties in 2012 have a significantly higher optical depth and may require some deceleration of the expansion rate.  This leaves the scenario developed by \\cite{Kochanek2011} for SN~2008S and the NGC~300-OT, where the absorption is due to a dense pre-existing wind. In this scenario, the progenitors were obscured by a dense, dusty wind where the dust is destroyed by the shock breakout luminosity from the explosive transient so as to leave the transient with little obscuration at peak. The densities are so high, however, that the dust can then reform and re-obscure the system, consistent with the early dust formation suggested by the spectra and photometry.  In the SN~2008S scenario, the material already exists at these distances and so dust can begin (re-)forming very rapidly.  If the luminosity is then powered by a combination of any surviving star and an expanding shock wave, the optical depth to the shock front is only dropping as $\\tau_V \\propto 1/t$. The resulting drop in the optical depth by only a factor of $4.9$ is far more compatible with the observations, as shown in Fig.~\\ref{fig:params}.  Arguably, the shock should also be slowing as it expands through the dense wind. Self-similar solutions (e.g. \\cite{Chevalier1983}) give expansion rates of $R \\propto t^s$ and $v_s \\propto t^{s-1}$ where $s=(n-3)/(n-2)$ for ejecta with an effective density profile $\\rho \\propto R^{-n}$ expanding into a $\\rho \\propto R^{-2}$ wind. The typical approximations for the structure of the ejecta are $n=7$ ($s=0.8$) and $n=12$ ($s=0.1$), leading to a modest slowing with time. In Fig.~\\ref{fig:params} we show the scaling of the conditions in 2004 to 2012 for this shock scenario ($v_s \\propto t^{-0.2}$, $R \\propto t^{0.8}$, $\\tau \\propto t^{-0.8}$).  If a significant part of the present day luminosity is driven by X-rays from an expanding shock, inferences about the temperature of the illuminating source and the exact value of the optical depth become somewhat problematic. In the shock scenario, the X-rays are absorbed by the dense gas and remitted as a complex, non-thermal (mainly) emission line spectrum, which is then absorbed and reradiated by the dust. While the shock emission models of \\cite{Allen2008} do not extend to the regimes considered here, they generally have relatively low near-IR emission compared to optical/UV emission and so would resemble a hot star model.  The estimates of the optical depth, which are driven by the near-IR detections, should also be viewed as being only logarithmically correct in this scenario.   The overall luminosity, dust temperature and dust radius, which are largely determined by the mid-IR emissions from the dust, should still be accurate.  If we fit the expansion of the dust radius as a power law, $R \\propto t^x$, we find $x=0.67$ ($0.26 < x < 1.05$, $1\\sigma$) for the graphitic models and $x=0.54$ ($0.29 < x < 0.80$, $1\\sigma$) for the silicate models.  The silicate models are consistent with linear expansion only at $2\\sigma$. Similarly, if we fit the evolution of the optical depth as a power law, $\\tau_V \\propto t^y$, we find $y=-1.27$ ($-1.69 < y < -0.90$, $1\\sigma$) and $y=-0.79$ ($-1.29 < y < -0.47$, $1\\sigma$) for the graphitic and silicate models, respectively.  The $\\tau_V \\propto 1/t^2$ scaling of expanding shells is ruled out at roughly $2\\sigma$.  We can also fit the evolution of the optical depth as a power law in radius, $\\tau_V \\propto R^z$, finding $z=-1.75$ ($-3.90 < z < -1.01$) and $z=-1.52$ ($-3.46 < z < -0.69$) for the two models.  While this appears to be consistent with the $\\tau_V \\propto R^{-2}$ scaling for an expanding shell, there is a strong covariance between this exponent and the expansion rate of the shell -- a solution with $\\tau_V \\propto R^{-2}$ must also roughly have $R \\propto t^{1/2}$.  Essentially, for the same change in optical depth there is considerable uncertainty in the radius because we lack a complete dust SED in 2012 for constraining the dust temperature. In the two-dimensional probability distribution of the $x$ and $z$ exponents, the expanding shell solution with $x=1$ and $z=-2$ is still ruled out at roughly $2\\sigma$.  The evolution is consistent with the $\\tau_V \\propto 1/R$ and $R \\propto t^{0.8-1.0}$ evolution of the expanding shock model.  A shock moving through a wind at velocity $v_s$ has a characteristic X-ray energy of \\begin{equation} E_s = { 3 \\mu \\over 16 } m_p v_s^2 = 1.2 \\left( { v_s \\over 1000~\\hbox{km/s}} \\right)^2~\\hbox{keV} \\end{equation} where the mean molecular weight is $\\mu=0.6$, and produces luminosity \\begin{equation} L_s =  { \\epsilon \\over 2 } { \\dot{M} \\over v_w } v_s^3 \\end{equation} where $\\epsilon \\leq 1$ is the efficiency with which the shock energy is radiated as X-rays (e.g. \\citealt{Chevalier1982}, \\citealt{Chugai1992}, \\citealt{Chugai1994}, \\citealt{Chevalier1994}).  The radius of the shock $ R=v_s t =3.1 \\times 10^{16} (v_s/1000~\\hbox{km/s})$~cm is simply related to its velocity, where $t=3566$~days  is the elapsed time from the transient to the X-ray observation.  This means that the hydrogen column density outside the shock radius is \\begin{eqnarray} N_H &= &{ \\dot{M} \\over 4 \\pi v_w  R (1.4m_p) } \\\\ &= &7.5 \\times 10^{21} \\left( { \\dot{M} \\over 10^{-4} M_\\odot/\\hbox{year} }\\right) \\left( { 10~\\hbox{km/s} \\over v_w} \\right) \\left( { 1000~\\hbox{km/s} \\over v_s }\\right)~\\hbox{cm}^2 \\nonumber \\end{eqnarray} Combining these three equations with the PIMMS models for absorbed thermal bremsstrahlung, we can determine the range of shock velocities and wind density parameters ($\\dot{M}/v_w$) that would violate the Swift limits, as shown in Fig.~\\ref{fig:params}.  If the shock velocity is as high as the $v_s = 893$~km/s FWHM reported by \\cite{Smith2011}, then there are allowed solutions for low wind densities (where the X-ray luminosity is low but little absorbed), and high wind densities (where the X-ray luminosity is high but heavily absorbed).  If the shock velocity $v_s < 600$~km/s, then the X-ray emission limits are satisfied for any wind density.  The parameters derived from the photometric fits broadly satisfy these limits on the X-ray emission. Half the graphitic solutions violate the limit in Fig.~\\ref{fig:params}, but the limit as drawn is for 100\\% conversion of the shock energy into soft X-rays and a $1\\sigma$ detection threshold, which is a rather optimistic representation of the detection threshold.  The basic picture of Fig.~\\ref{fig:params} is little changed if we use the larger distances to NGC~4242, with all solutions and limits simply shifting to higher velocities.  While this increases the shock luminosity, the net effect is limited because the X-ray luminosity limits from the observations become correspondingly weaker."
    },
    "1208/1208.5179.txt": {
        "abstract": "{We review the mechanisms via which an external magnetic field can affect the ground state of cold and dense quark matter. In the absence of a magnetic field, at asymptotically high densities, cold quark matter is in the Color-Flavor-Locked (CFL) phase of color superconductivity characterized by three scales:  the superconducting gap, the gluon Meissner mass, and the baryonic chemical potential. When an applied magnetic field becomes comparable with each of these scales, new phases and/or condensates may emerge.  They include the magnetic CFL (MCFL) phase that becomes relevant for fields of the order of the gap scale; the paramagnetic CFL, important when the field is of the order of the Meissner mass, and a spin-one condensate associated to the magnetic moment of the Cooper pairs, significant at fields of the order of the chemical potential.  We discuss the equation of state (EoS) of MCFL matter for a large range of field values and consider possible applications of the magnetic effects on dense quark matter to the astrophysics of compact stars.} ",
        "introduction": "\\label{Intro}  Contrary to what our na\\\"{\\i}ve intuition might indicate, a magnetic field does not need to be of the order of the baryon chemical potential to produce a noticeable effect in a color superconductor. As discuss in \\cite{SpinoneCFL,phases}, a color superconductor can be characterized by various scales and different physics can occur at field strengths comparable to each of them. Specifically,  for the so-called Color-Flavor-Locked (CFL) phase, the superconducting gap, the Meissner mass of the charged gluons, and the baryon chemical potential define three scales that determine the values of the magnetic field needed to produce different effects. Thus, the presence of sufficiently strong fields can modify the properties of the dense-matter phase which in turn might lead to observable signatures.  In this paper we review the status of our current knowledge about the magnetic field effects on color superconductivity (CS) at asymptotically high densities and discuss possible consequences of these effects for the physics of compact stars. ",
        "conclusions": ""
    },
    "1208/1208.5127_arXiv.txt": {
        "abstract": "{} {Recently, the study of globular cluster (GC) color-magnitude diagrams (CMDs) has shown that some of them harbor multiple populations with different chemical compositions and/or ages. In the first case, the most common candidate is a spread in the initial helium abundance, but this quantity is difficult to determine spectroscopically due to the fact that helium absorption lines are not present in cooler stars, whereas for hotter GC stars gravitational settling of helium becomes important. As a consequence, indirect methods to determine the initial helium abundance among populations are necessary. For that reason, in this series of papers, we investigate the effects of a helium enrichment in populations covering the range of GC metallicities. } {In this first paper, we present the theoretical evolutionary tracks, isochrones, and zero-age horizontal branch (ZAHB) loci calculated with the Princeton-Goddard-PUC (PGPUC) stellar evolutionary code, which has been updated with the most recent input physics and compared with other theoretical databases. The chemical composition grid covers 9 metallicities ranging from Z=1.60$\\times 10^{-4}$ to 1.57$\\times 10^{-2}$ (-2.25$\\lesssim$[Fe/H]$\\lesssim$-0.25), 7 helium abundances from Y=0.230 to 0.370, and an alpha-element enhancement of [$\\alpha$/Fe]=0.3. } {The effects of different helium abundances that can be observed in isochrones are: splits in the main sequence (MS), differences in the luminosity ($L$) and effective temperature ($T_{\\rm eff}$) of the turn off point, splits in the sub giant branch being more prominent for lower ages or higher metallicities, splits in the lower red giant branch (RGB) being more prominent for higher ages or higher metallicities, differences in $L$ of the RGB bump (with small changes in $T_{\\rm eff}$), and differences in $L$ at the RGB tip. At the ZAHB, when Y is increased there is an increase (decrease) of $L$ for low (high) $T_{\\rm eff}$, which is affected in different degrees depending on the age of the GC being studied. Finally, the ZAHB morphology distribution depending on the age explains how for higher GC metallicities a population with higher helium abundance could be hidden at the red ZAHB locus. } {} ",
        "introduction": "\\label{intro}  Globular clusters (GCs) are objects orbiting around the Milky Way and other galaxies and containing between $\\sim$10$^4$ and $\\sim$10$^7$ stars. Nowadays, $\\sim$160 GCs are known in the Milky Way, with metallicities $[M/H] = \\log (Z/X) -\\log (Z/X)_\\odot$ between $\\sim$-2.3 and $\\sim$-0.25 \\citep[][revision 2010]{Harris1996} and with ages ranging from $\\sim$ 6.5 to $\\sim$13 Gyr \\citep{Aparicio_etal2001, Chaboyer2001, MarinFranch_etal2009, Dotter_etal2010, MoniBidin_etal2011}.  For several decades these objects have been considered as an excellent laboratory for the study of simple stellar populations (SSPs) because it was believed that all stars inside them were formed from a homogeneous cloud at the same time. However, it was recognized early that their horizontal branch (HB) morphologies could not be simply parametrized.  While the metallicity of a GC is the first parameter defining its HB morphology \\citep{Sandage_Wallerstein1960}, empirical findings show clearly that it cannot be the only parameter at work. In fact, for many years we have been well aware of several GCs characterized by quite different HB morphologies in spite of their quite similar metallicity. For this reason, several {\\em second parameters} have been invoked to solve this conundrum, where, of course, the most direct, and also the first invoked, second parameter was the age \\citep{Sandage_Wallerstein1960, Rood_Iben1968}. This problem has been widely debated until the present date, but without reaching a consensus \\citep[e.g.][ and references therein]{Catelan_etal1993, Lee_etal1994, Stetson_etal1996, Sarajedini_etal1997, Ferraro_etal1997c, Catelan2009b, Dotter_etal2010, Gratton_etal2010}.  Another early invoked candidate was the initial helium abundance \\citep{van_den_Bergh1965, van_den_Bergh1967, Sandage_Wildey1967, Hartwick1968}, but the major difficulty with this element is its extremely difficult measurement, which requires high temperatures in order to be detected in stars. Moreover, blue HB stars hotter than $\\sim$11\\,500 K as well as stars on the so-called extreme HB suffer from metal levitation and helium settling \\citep[e.g.,][]{Grundahl_etal1999,Behr2003}, leaving only a narrow band between $\\sim$11\\,500 and $\\sim$8\\,000 K for detecting the initial helium abundance in HB stars. Even so, its measurement remains difficult \\citep[see][]{Villanova_etal2009, Villanova_etal2012}. Even though helium was one of the first candidates, this explanation of the second-parameter effect became less popular after the Milky Way formation scenario of \\citet{Searle_Zinn1978}, where the age is assumed to be the second parameter. In addition, differences in helium among GCs also seemed less likely, due to the increasing evidence that GCs have similar helium abundances \\citep[e.g.,][]{Salaris_etal2004}.  However, after the observational evidence of multiple main sequences (MSs) in GCs \\citep{Norris2004, Piotto_etal2005, DAntona_etal2005, Piotto_etal2007}, the hypothesis that helium is the second parameter has regained strength, but in this case the difference in the initial helium abundance is mostly among stars inside the same GC \\citep[][among several others]{DAntona_etal2002, DAntona_etal2005, Norris2004, Lee_etal2005, Piotto_etal2005, Piotto_etal2007}. In fact, it seems that the combination of age as the second parameter and a helium spread among stars as the ``third parameter'' is a promising attempt \\citep{Gratton_etal2010}.  It is worth noting that, although the existence of a correlation between the chemical abundance anomalies (such as light-element anticorrelations) and the HB morphology has been suggested for a long time \\citep{Norris1981, Smith_Norris1983, Catelan_deFreitasPacheco1995}, only recently has direct evidence of this correlation been obtained, as in the case of the GC M4 \\citep[][]{Marino_etal2011, Villanova_etal2012}.  As can be noted, the hypothesis that GCs are SSPs is being ruled out, due to the chemical inhomogeneities among stars of the same GC that have been (directly or indirectly) detected and associated to different star formation episodes in the same GC \\citep[][and references therein]{Decressin_etal2007b, DErcole_etal2008, Carretta_etal2010, Conroy_Spergel2011, Bekki2011, Valcarce_Catelan2011}.  In addition to the presence of multiple populations in GCs, another development makes the availability of evolutionary tracks for different helium abundances highly necessary: as pointed out by \\citet{Catelan_deFreitasPacheco1996} and \\citet{Catelan2009}, and also emphasized recently by \\citet{Nataf_Gould2012}, \\citet{Nataf_Udalski2011}, and \\citet{Nataf_etal2011,Nataf_etal2011a}, stellar populations with different enrichment histories may have different helium enrichment ``laws\" (i.e., different Y-Z relations).  In this article, we show the effects of the helium enrichment in isochrones and zero-age HB (ZAHB) loci in the theoretical Hertzsprung-Russell (HR) diagram. First in section \\ref{PGPUCsec}, we present the Princeton-Goddard-PUC (PGPUC) stellar evolution code (SEC) which we used to create a set of evolutionary tracks, isochrones and ZAHB loci for different metallicities and helium abundances (section \\ref{PGPUCdatabase}). Then, we show the effects of the helium enrichment (section \\ref{EffectsHe}) and end with the conclusions (section \\ref{conclu})\\footnote{All of the models described in this paper can be downloaded from the ``PGPUC Online\" dedicated website, as described in Appendix \\ref{AppPGPUConline}.}. ",
        "conclusions": "\\label{conclu}  In this work, we studied the effects that must be observed in the theoretical $\\log (L/L_\\odot)-\\log (T_{\\rm eff})$ plane for isochrones and ZAHB loci with different initial helium abundances and fixed metallicity, which is the most common hypothesis to explain the multiple populations observed in the color-magnitude diagrams in some GCs. Our study is based on stellar evolutionary calculations for different evolutionary phases in the life of a low-mass star.  First, we presented the PGPUC SEC that was updated with the most recent physical ingredients available in the literature, including opacities, both radiative and conductive \\citep{Iglesias_Rogers1996, Ferguson_etal2005, Cassisi_etal2007}; thermonuclear reaction rates \\citep{Angulo_etal1999, Kunz_etal2002, Formicola_etal2004, Imbriani_etal2005}; equation of state \\citep{Irwin2007}; mass loss \\citep{Schroder_Cuntz2005}; and boundary conditions \\citep{Catelan2007}. Even though the neutrino energy loss rate was not updated, we recommend the use of the \\citet{Kantor_Gusakov2007} interpolated tables for extreme conditions (high densities), which however are not achieved for the stars in our study.  Comparison between PGPUC SEC and other theoretical databases shows similar results. In fact, PGPUC SEC evolutionary tracks and isochrones are in excellent agreement with those produced using the PADOVA, GARSTEC and (specially) BaSTI codes.  We create a large set of evolutionary tracks for a wide range of metallicities and helium abundances for a GS98 chemical composition and [$\\alpha$/Fe]=+0.3, which are available in the PGPUC Online web page. Then theoretical isochrones and ZAHB loci were used to show the effects of different helium abundances at different evolutionary phases along the HR diagram, where in isochrones:  \\begin{itemize} \\item At the MS, the separation induced by Y increases for high Z.  \\item At the SGB, the separation induced by Y increases for high Z, but it decreases with the age.  \\item At the base of the RGB, the difference in $T_{\\rm eff}$ induced by Y increases for high Z and also increase with the age. \\end{itemize}  \\noindent while at the ZAHB loci:  \\begin{itemize} \\item When Y increases, the $L$ of the cooler HB stars increases, but the $L$ of the hotter HB stars decreases.  \\item When the age increases, the $L$ of the hotter HB stars increase due to the increase of the $M_{cHe}$.  \\item When Z increases, there is an increase in the separation in $L$ due to a difference in Y at the cooler HB stars, but the luminosity difference induced by the age decreases. \\end{itemize}  We emphasized the importance of the age in determining the ZAHB locus \\citep[as pointed out by][and references in Sect.\\ref{ZAHBage}]{Caputo_deglInnocenti1995}, and also the location of stars when they arrive to the ZAHB loci depending on the age. In the next papers of this series we will study these effects in different filters for comparing with GCs data, and we will also create a new set of evolutionary tracks with different CNO ratios.  \\subsection*"
    },
    "1208/1208.0702_arXiv.txt": {
        "abstract": "{The knowledge of the present-day mass function of young clusters and the mass of their coolest substellar members is essential to clarify the brown dwarf formation mechanism, which still remains a matter of debate. } {We searched for isolated planetary-mass T-dwarfs in the $\\sim$3~Myr old Serpens Core cluster. } {We performed a deep imaging survey of the central part of this cluster using the WIRCam camera at the CFHT. Observations were performed through the narrow-band CH$_4$off and CH$_4$on filters, to identify young T-dwarfs from their 1.6$\\mu$m methane absorption bands, and the broad-band $JHK_S$ filters, to better characterize the selected candidates. We complemented our WIRCam photometry with optical imaging data from MegaCam at CFHT and Suprime-Cam at the Subaru telescope and mid-infrared flux measurements from the Spitzer  ``core to disk'' (c2d) Legacy Survey.} {We report four faint T-dwarf candidates in the direction of the Serpens Core with CH$_4$on$-$CH$_4$off above 0.2~mag, estimated visual extinction in the range 1-9~mag and spectral type in the range T1-T5 based on their dereddened CH$_4$on$-$CH$_4$off colors. Comparisons with T-dwarf spectral models and optical to mid-infrared color-color and color-magnitude diagrams, indicate that two of our candidates (ID~1 and 2) are background contaminants (most likely heavily reddened low-redshift quasars). The properties of the other two candidates (ID~3 and 4) are consistent with them being young members of the Serpens Core cluster, although our analysis can not be considered conclusive. In particular, ID~3 may also be a foreground T-dwarf. It is detected by the Spitzer c2d survey but only flux upper limits are available above 5.8~$\\mu$m and, hence, we can not assess the presence of a possible disk around this object. However,  it presents some similarities with other young T-dwarf candidates (S~Ori~70 in the $\\sigma$~Orionis cluster and CFHT\\_J0344+3206 in the direction of IC~348). If ID~3 and 4 belong to Serpens, they  would have a mass of a few Jupiter masses and would be amongst the youngest, lowest mass objects detected in a star-forming region so far.} ",
        "introduction": "Understanding the brown dwarf (BD) formation mechanism is a key point to assess what determines and which is the minimum mass for star formation. Although two decades have passed since the first unambiguous observations of BDs \\citep{Reb95,Nak95}, this issue is still under debate. According to the star formation theory, the minimum mass for star formation is set by the so called mass fragmentation limit, i.e. the mass at which one object cannot contract further because the radiated heat prevents it to collapse further. This limit is of the order of a 5-10M$_J$ \\citep{Low76,Ree76} and may be lower ($\\sim$1~M$_J$) in the presence of magnetic fields \\citep{Bos01,Boy05}. Observationally, several young star forming regions have been probed to faint magnitudes at both optical and near-infrared wavelengths to look for a possible cut-off at the low-mass end of the mass function and the results significantly vary from region to region. A possible cut-off below $\\sim$6~M$_J$ up to 20~M$_J$ have been reported in NGC~1333 \\citep[e.g.,][]{Sch09,Sch12}. Moreover, there is no universal agreement on the behavior of the mass function over this cut-off. A search for T-dwarfs in IC~348 by \\citet{Bur09} suggests that an extrapolation of the log-normal mass function \\citep{Cha03} may hold in the substellar and planetary-mass regimes. A similar study in the Upper Sco association favors a turn-down in the mass function below 10~M$_J$ \\citep{Lod11}. In the Pleiades \\citet{Cas07} reported several L/T-type candidates with masses as low as 10~M$_J$, implying that the slope of the mass function in this mass regime agrees, within the uncertainties, with the values inferred from earlier studies at higher masses \\citep{Dob02,Mor03,Lod07,Cas07}. Finally, \\citet{Kir12} find tantalizing hints that the number of BDs in the field continues to rise from late-T to early-Y.  To address these questions, a large observing key-program (P.I. J. Bouvier) was conducted at the Canada France Hawaii Telescope (CFHT) aimed at characterizing the stellar population of nearby star forming regions using the Wide Field IR Camera (WIRCam). One of the main objectives of this program is to look for the lowest mass members of these clusters and investigate the cut-off at the low mass end of the mass function, using a technique based on narrow-band methane imaging \\citep[e.g.,][]{Bur09}.  In this paper, we report the results of this observing program in the Serpens Core cluster, specifically focusing on the methane filters imaging to search for young T-dwarfs candidates. The Serpens Core region is an example of a very young, deeply embedded (A$_V \\approx$40~mag) cluster, containing a high percentage of protostars \\citep{Dav99,Tes00,Kaa04,Eir06,Har07,Win07}. With an age between 2 and $\\sim$6~Myr \\citep{Win09,Oli09} and a distance estimated between 260$\\pm$37~pc \\citep{Str03} and 415$\\pm$5~pc \\citep{Dzi10}, the Serpens cloud core is one of the nearest regions of clustered star formation to the Sun. Therefore, it is an excellent candidate for study as it is close enough to both resolve the individual members and detect the lowest mass members to below the hydrogen-burning limit. Throughout this paper, we assume a typical age of 3~Myr \\citep{Win09} for the Serpens core cluster and the minimum distance of 260$\\pm$37~pc \\citep{Str03} \\footnote{This choice is appropriate because we use the distance to Serpens to estimate the lower limit to the interstellar extinction towards the cloud (see also Sect.~\\ref{sect_par}).}. The region is also particularly interesting for BD studies because only a few dedicated searches have been reported in the literature. The first young BD in Serpens was discovered by \\citet{Lod02} and the cluster substellar population count so far about 3 spectroscopically confirmed BDs  \\citep{Lod02,Shi11} and 40 BD candidates, a considerable fraction of which still surrounded by prominent accretion disks \\citep{Eir06}.  Canonical approaches on BD studies and related issues \\citep[see, e.g.,][]{Jay03} rely on the identification and spectral classification of BDs. However, optical/near-infrared (IR) spectroscopy \\citep[see, e.g.,][]{McL03,Geb02} has proved impossible for very faint distant BDs with current telescopes and very time-consuming for large BD samples. To overcome these limitations, new BD classification methods have been devised on the basis of near-IR broad-band \\citep[see, e.g.,][and references therein]{Leg10,Sch10} and narrow-band imaging that targets unique molecular features of L and T dwarfs, in particular, water and methane bands \\citep[see, e.g.,][]{Gor03,Mai04,Spe11}. These techniques are equivalent to extremely low-resolution spectroscopy and can be confidently applied for statistical purposes, e.g., to detect and classify \\emph{bona fide} BDs in large imaging surveys. Our search for T-dwarfs in Serpens is based on narrow-band methane (CH$_4$) imaging. We target, in particular, the CH$_4$ absorption band centered at 1.66~$\\mu$m, which forms in the atmosphere of objects cooler than $T_{eff} \\lesssim$1200~K \\citep[defyning the L/T dwarf boundary;][] {Bur06}. This technique has been successfully applied to a search for T-dwarfs in the young IC~348 cluster \\citep{Bur09}.   This paper is organized as follows. Section~\\ref{obs} describes the WIRCam observations in the Serpens Core and the catalog extraction procedure. In Sects.~\\ref{sel}-\\ref{sect_par}  we use this catalog to identify possible T-dwarfs on the basis of methane absorption bands and estimate their intrinsic properties (T$_{eff}$, A$_V$, etc.). In Sect.~\\ref{discussion} we investigate the nature of our T-dwarfs candidates using complementary near to mid-IR data and T-dwarf spectral models. The conclusions of this work are given in Sect.~\\ref{concl}. ",
        "conclusions": "}   We presented a deep methane imaging survey for planetary-mass T-dwarfs in the Serpens Core cluster. Our survey covers a field of view of 21.5$^\\prime \\times$21.5$^\\prime$ and about 127000 sources were detected in both the CH$_4$on and CH$_4$off filters.  We identified four potential T-dwarfs from their 1.6$\\mu$m methane absorption band and used complementary $riJHK_S$ broad-band imaging and mid-IR flux measurements from the {\\it Spitzer} c2d Survey to investigate their stellar and disk properties.  Two of our candidates (ID~1 and ID~2) suffer from significant interstellar extinction and have near-IR colors similar to mid-T dwarfs. However, they are too bright to be planetary-mass members of Serpens and do not present the typical IR excess emission expected for young objects surrounded by a disk. Our analysis indicates that their dereddened near-IR colors approach the values expected for QSOs. In particular, since they are too bright to be high-redshift QSOs, they could be heavily reddened low redshift QSOs.  Our candidate ID~3 does not present a significant interstellar extinction and its near-IR colors are more consistent with it being a field mid-T dwarf. It is brighter than field objects at the distance of Serpens and, hence, it may be a foreground T-dwarf. It is detected by the {\\it Spitzer} c2d survey but only flux upper limits are available for $\\lambda \\ge$5.8$\\mu$m and, hence, we can not assess the presence of a possible disk around this object. However, it is worth to keep this possibility open, because the position on the $J$ vs. $J-K_S$ and $i-J$ vs. $J$-[3.6] diagrams and the estimated spectral type of ID~3 are very similar to those of other young T-dwarf candidates, i.e. S~Ori~70 and  CFHT\\_J0344+3206 in the direction of IC~348.  Finally, ID~4 is significantly extincted, which exclude the possibility of it being a foreground object, and has dereddened optical and near-IR colors mainly consistent with those of early/mid T-dwarfs. It is brighter than field objects and, hence, it might be very young, as YSOs have larger radii and are brighter than older field objects of the same spectral class. Thus, ID~4 is a promising young T-dwarf candidate in our sample.  Considering the magnitudes of our candidates ID~3 and 4 (Table~\\ref{tab_phot}) and their T$_eff$ (Table~\\ref{tab_par}) based on the dereddened CH$_4$on$-$CH$_4$off colors, we estimated their masses assuming that they belong to the Serpens core cluster (i.e., age$\\sim$3~Myrs and distance in the range 260-415~pc) and using both the COND and DUSTY evolutionary models by \\citet{Cha00}. The estimated mass is in the range 2-4 Jupiter masses for both ID~3 and 4. Therefore, if they truly belong to the Serpens Core cluster, they would be amongst the youngest, lowest mass objects detected in a star-forming region so far.  Follow-up spectroscopy is required to confirm the spectral type and reddening  of our four T-dwarf candidates and, hence, draw firm conclusions about their nature. Additional deep imaging at mid-IR wavelengths is needed to clarify the presence of a possible disk around ID~3."
    },
    "1208/1208.3537_arXiv.txt": {
        "abstract": "Recent studies have indicated that large mass radial inflow rates are possible in observed galaxies with strong density wave patterns. Yet, numerical simulations have generally failed to account for such high rates. Here it is shown that the reason for the discrepancy is the treatment of ``softening'', the artificial parameter inserted by numerical simulators into the formula for gravitational potential, to control the magnitude of relaxation in simulations with small numbers of particles compared to a real galaxy. Excess softening reduces the collective effects underlying the significant secular evolution inferred for physical galaxies. Less softening, coupled with an increase in number of particles, allow the N-body simulations to reveal significant morphological transformation of galaxies over a Hubble time. ",
        "introduction": "The prospect of significant transformation of the Hubble type of a galaxy during its lifetime has important implications on our understanding of the origins of the Hubble sequence and on broader issues in cosmology.  Despite the coming of age of this field of research, there is as of now no consensus on how important this process is, and through what dynamical mechanism this transformation is mainly achieved (see the review of Kormendy \\& Kennicutt 2004).  Zhang (1996, 1998, 1999) proposed and demonstrated a collective dissipation process in galaxies that could serve as the underlying dynamical mechanism for secular evolution of galaxies along the entire Hubble sequence, by incorporating the stellar component in the mass inflow process rather than just the interstellar medium as proposed in the pseudo-bulge formation scenario (Kormendy 1979). Despite the confirmation of the theoretically derived mass flow rates in the accompanying N-body simulations, which established the viability of the analytical approach, the simulated mass flow rates were small, far from being adequate to lead to significant morphological transformation of a galaxy within a Hubble time.  This small mass flow rate was attributed to the small amplitudes of the spiral patterns formed in these simulations, with the state-of-the-art in simulating spontaneously formed modes far from being adequate to produce the extremely nonlinear density wave modes observed in physical galaxies.  In order to circumvent this difficulty, Zhang \\& Buta (2007; 2012) used near-and-mid-infrared images of galaxies, coupled with the analytical rate equations derived in Zhang (1996, 1998, 1999) to calculate the torques and mass flow rates in these physical galaxies, and found that mass flow rates from a few $M_{\\sun}$ per year to over one hundred $M_{\\sun}$ per year were obtained, depending on the Hubble types of galaxies and their interacting environment, which together determined the density wave amplitudes and pitch angles in these galaxies. These mass flow rates for real galaxies far exceed those obtained from the past N-body simulations, partly as a result of the fact that the mass flow rate is proportional to the effective wave amplitude squared (Zhang 1998), and the amplitudes in physical galaxies are oftentimes a factor of 3-8 times higher than obtained in N-body simulations, implying a difference in evolution rate of a factor of 10 - 100 between physical and previously-simulated galaxies.  Despite the confirmation of the importance of secular evolution process in physical galaxies in works such as Zhang \\& Buta (2007,2012), one could not help but wonder what exactly were the factors that have prohibited the simulated disk galaxies from obtaining the level of mass flow rates and the magnitude of wave amplitudes in physical galaxies.  It is to this question that the result of the current Letter partly addresses.  Incidentally, the main discovery of this work was made in a totally serendipitous fashion -- i.e. the result was stumbled upon while the author was searching in a different direction for the cures of the low mass flow rates in N-body simulations\\footnote{The author used to think that the low mass flow rates were mainly due to the two-dimensional (2D) nature of her simulations, and once the bulge and halo were allowed to be active, the mass flow rates will become higher.  This turned out to be true, but only moderately true: i.e., having about 15\\% active bulge and halo in 2D simulations increases the mass inflow rate by about 10\\%, but having still larger active bulge and halo fraction leads rather to outflows of mass over the main parts of the disk (likely due to the lowering of Q and the emergence of transient local instabilities), contrary to what is needed for bulge building. Note that the above conclusion about the fraction of allowable active bulge/halo mass applies only in the 2D cases where these tests had been made.  In three dimensional (3D) studies the entire bulge and halo mass of course can be made active with proper pressure and rotational support.}). The lesson learned is that one has to be cautious in taking the results of N-body simulations literally without questioning sacred ``rules of thumb'' and accepted wisdom, especially in circumstances where new physical processes are encountered -- in our situation, that is when the collective effects become important. ",
        "conclusions": "The low mass inflow rates found in the past N-body simulations of disk galaxies had cast doubt on the effectiveness of secular evolution as an important dynamical process for the morphological transformation of galaxies along the Hubble sequence. In this work it is shown that such low numerical mass inflow rates are chiefly the result of the artificial ``softening'' of gravity which is a common practice in galaxy simulations to avoid rapid relaxation due to the small number of particles employed.  By decreasing the amount of softening and simultaneously increasing the number of simulation particles, realistic level of mass inflow comparable to those inferred from observed galaxy images through torque calculations are achieved in N-body simulations of disk galaxies, which reinforce the importance of secular evolution as an extremely relevant process in transforming the morphologies of galaxies during the past Hubble time."
    },
    "1208/1208.3647_arXiv.txt": {
        "abstract": "We study vacuum quantum fluctuations of simple Nambu-Goldstone bosons - derivatively coupled single scalar-field theories possessing shift-symmetry in field space. We argue that quantum fluctuations of the interacting field can be drastically suppressed with respect to the free-field case. Moreover, the power-spectrum of these fluctuations can soften to become \\textit{red} for sufficiently small scales. In \\textit{quasiclassical} approximation, we demonstrate that this suppression can only occur for those theories that admit such \\textit{classical static} backgrounds around which small perturbations propagate faster than light. Thus, a \\textit{quasiclassical softening} of quantum fluctuations is only possible for theories which \\textit{classicalize} instead of having a usual Lorentz invariant and local Wilsonian UV- completion. We illustrate our analysis by estimating the quantum fluctuations for the DBI-like theories. ",
        "introduction": "Quantum Mechanics (QM) of a finite number of degrees of freedom is a paradigm which should a priory be applicable to arbitrary Hamiltonian systems including strongly coupled models. In particular, QM describes systems with noncanonical and even highly nonlinear kinetic terms as well as it works for those with the highly nonlinear potentials. However, currently, in the case of QM for a continuum number of degrees of freedom (i.e. for quantum field theory (QFT)) all models well defined at all scales (renormalizable) are just quadratic in canonical momentum. For QFT a noncanonical structure of the kinetic term usually implies the existence of a strongly coupled regime where the usual perturbative renormalization procedure breaks down and where one has to use non-perturbative methods, e.g. lattice computation. However, lattices of different size correspond to the approximation of QFT by finite dimensional QM systems with different number of degrees of freedom and the convergence of this procedure is not obvious.  It is a common belief that Nature is such that it can be described in terms of fields which are different in the weak and strong coupling regimes. The best examples for this behavior are given within the Standard Model by QCD and Higgs field in the Electroweak sector. One tends to think that in UV the number of degrees of freedom - types of particles - should increase, see e.g. most recent works \\cite{Komargodski:2011vj,Luty:2012ww}. New particles should be integrated in to provide a UV complete theory - this is the basis of the Wilsonian UV-completion. This is the case for String Theory where in the UV there is a continuum of fields. In this regard natural questions arise: is continuum QM - QFT -infinitely more restrictive than QM for finite number of degrees of freedom? Can noncanonical or perturbatively non-renormalizable QFTs make sense in a continuum limit i.e. for \\textit{all scales}? If it is the case, can one understand the mechanism used by Nature in these theories in terms of some weakly coupled fields / particles? These questions could be crucial for understanding of Quantum Gravity which has a noncanonical Hamiltonian with the strong-coupling scale given by the Planck mass.  Recently it was proposed \\cite{Dvali:2010bf,Dvali:2010jz,Dvali:2010ns,Dvali:2010ue} that gravity may be self-UV-completed via \\textit{classicalization} - the softening of quantum fluctuations i.e. loops in Feynman diagrams at the transplanckian transferred momenta, because of the formation of intermediate Black Holes\\footnote{For an earlier similar discussion of the quasiclassical high-energy scattering in gravity, see e.g. \\cite{Banks:1999gd}.}. \\textit{Classicalization} is caused by the high level of nonlinearity and corresponding self-sourcing. Moreover, it was suggested that this mechanism may also work for quite generic class of the Nambu-Goldstone bosons \\cite{Dvali:2011th,Dvali:2010ns}\\footnote{For criticism of this interesting proposal, see e.g. Refs. \\cite{Akhoury:2011en,Kovner:2012yi}} and even more general scalar fields including nonlinear sigma models \\cite{Percacci:2012mx}, where the role of Black Holes is played by \\textit{classicalons} - extended and long lived classical configurations. However, this phenomenon was argued \\cite{Dvali:2011nj,Dvali:2012zc} to occur for those models for which the usual local and Lorentz-invariant Wilsonian UV-completion is absent, because they allow for a superluminal propagation of perturbations around nontrivial backgrounds \\cite{Adams:2006sv} \\footnote{Here it is interesting to note that already in 1952 Heisenberg argued that derivative self-interaction is important for the correct description of high multiplicity in scattering in strongly interacting theories. See \\cite{Heisenberg:1952zz} where the usual subluminal DBI is used to analyse this problem. % }.  In this paper we analyse the selfconsistency of \\textit{classicalization} for general noncanonical Nambu-Goldstone bosons in $d+1$ spacetime dimensions. We did not restrict our attention to $d=3$, because in lower spacial dimensions there are much higher chances to find theories for which our estimations can be checked by numerical simulations or by other non-perturbative methods. We use the Heisenberg uncertainty relation for canonically conjugated field and momentum operators averaged on a given scale to estimate vacuum quantum fluctuations of the field on this scale. In our analysis we follow the ideas of \\textit{classicalization} and work in the \\textit{quasiclassical} approximation - neglecting higher-order correlators. On top of that we use analytic continuation which is a standard tool for interacting QFT with the energy functional bounded from below. We show that vacuum quantum fluctuations, $\\delta\\phi$, can be suppressed with respect to the free-field case only if the theory admits \\textit{static classical} backgrounds which are either i) absolutely unstable (have gradient instability) or ii) such that perturbations around them propagate faster than light. Option i) is definitely even more worrisome than ii) and should be excluded. Moreover, if the suppression is such that the spectrum of perturbations is \\textit{red} in the UV, then the speed of perturbations $v^{2}>d\\geq1$. We illustrate our finding by analysis of quantum fluctuations for Dirac-Born-Infeld (DBI) like theories. ",
        "conclusions": "We have demonstrated that the quasiclassical suppression of quantum fluctuations of Nambu-Goldstone bosons is only possible for theories which either possess catastrophically unstable backgrounds or allow for superluminality. This finding supports the analysis performed before in \\cite{Dvali:2011nj,Dvali:2012zc}.  It would be very interesting to study whether these results can be generalised to other systems and quantum states."
    },
    "1208/1208.3471_arXiv.txt": {
        "abstract": "We introduce  a new method for the identification of galaxy systems in redshift surveys based on the halo model. This method is a modified version of the K-means identification algorithm developed by \\cite{yang:04}. We have calibrated and tested our algorithms using mock catalogs generated using the Millennium simulations \\citep{springel:05} and applied them to the NYU-DR7 galaxy catalog (based on the SDSS datasets).  Using this local sample of groups and clusters of galaxies we have measured the effect of gravitational redshift produced by their host dark matter haloes. Our results shows radial velocity decrements consistent with general relativity predictions and previous measurements by \\cite{whh:11} in cluster of galaxies. ",
        "introduction": "In a recent article \\citep{whh:11} (WHH) analyze the pattern of spectroscopic redshifts for galaxies in 7800 clusters from the 7th data release of the Sloan Digital Survey (SDSS). Motion of a galaxy in a cluster generates a red or blueshift of its spectral lines; the equivalent velocities are typically on the order of 600 kms/s. WHH show that it is possible to disentangle the superimposed gravitational redshift predicted by the General relativity that corresponds to a radial velocity differences of the order 10 km/s. In their analysis the authors used a stacking data from several clusters and determined the shift of the redshift distribution's centroid with  growing radial coordinate. The positions and redshifts of cluster galaxies are derived from a Gaussian Mixture Brightest Cluster Galaxy cluster catalogue \\citep{hao:10}, one of the largest samples of galaxy clusters assembled on the basis of the SDSS-DR7. \\\\ In spite of the fact that massive clusters are expected to have the largest gravitational redshift effects, these systems are also affected by large galaxy peculiar velocities and substructure. Kim and Croft \\citep{kc:04} suggested  that it is possible to overcome partially these difficulties by averaging over many clusters and groups of relatively low mass.\\\\ The existence of volume complete samples of galaxies with redshifts measured in the nearby universe as in the SDSS main galaxy sample \\citep{abazajian:09}; lead us to explore the gravitational redshift effects using systems of galaxies in the nearby universe. Several groups have derived  catalogs of systems of galaxies applying  different techniques (Friends of friends FoF, Matched Filters etc.) to large data set such as the 2dF and the SDSS.\\\\ In order to make a meaningful comparison between observation and  theory, group definitions must have a reliable three dimensional counterpart. From the point  of view of  current theory of galaxy  formation,  the most direct route  for defining  galaxy groups is  dark matter halos. More important for this work is the novel hybrid technique introduced by \\cite{yang:04},  that starts with cluster candidates selected by a classical FoF  algorithm and a selected sample of bright isolated galaxies. Other satellite galaxies are classified according a Maximum Likelihood criterion. Such method employs a model of the redshift space distribution of the galaxies (as function of system mass) as filter. A convolution of the model with the galaxy distribution allows to decide the membership of each galaxy to the proposed systems (provided by a suitable and calibrated FoF method) using an iterative procedure. The method relies in an assumed mass to light relation, which is a more reliable  method for compute the system mass that those based on the computation of the velocity dispersion and virial radii from galaxy positions.\\\\ \\cite{yang:04} method is a basic version of the well known K-means, a clustering algorithm where points are assigned to exactly one cluster and all points assigned to a cluster are equals in that system. Yang have carefully tested the performance of their group finder. The method was applied to the 2dFGRS  and compared with those extracted from detailed mock galaxy  redshift  surveys.  More recently \\cite{yang:07} applied this method to the fourth release of the SDSS survey in order to study the dependence of color, stellar mass, star formation rate and morphology on halo mass.\\\\ In order to measure the gravitational redshift caused by dark matter haloes in systems of galaxies, we present in Section 2 modifications of the algorithm consisting in the introduction of a soft degree of assignment  of a galaxy to each cluster.  We have tested and calibrated our modified method in Section 3 with an extensive use of a galaxy mock catalog building up on the results of the Millenium simulations which mimics the  SDSS-DR7 galaxy catalog. In section 4  we use a selected sample of clusters and groups of galaxies from the SDSS-NYU-DR7 datasets \\citep{blanton:nyu} to compute the gravitational redshift produced by its dark matter haloes.  This is followed by a brief summary. We adopt for our study a flat concordance $\\Lambda CDM$  cosmology \\citep{komatsu:11} with $H_0= 72 \\kms \\mpcnoh, \\Omega_M=0.26, \\Omega_\\Lambda=0.74$,  where appropiate we define $h=H/100/\\kms/\\mpcnoh$.\\\\ ",
        "conclusions": "\\begin{itemize} \\item An efficient halo finding algorithm has been developed and tested using mock catalogs derived from numerical simulations with semianalytic galaxy formation methods.  The performance of the method is solid and efficiently provides complete, almost bias free galaxy system catalogs based on galaxy redshift surveys. \\item The new method was used to compile samples of groups and clusters of galaxies in the SDSS-NYU-DR7. The clusters of galaxies sample was used to reproduce the results found by WHH in clusters of galaxies. Following the suggestion by \\citep{kc:04}, we are able to detect the gravitational redshift in the sample of groups of galaxies. \\item Our measurements of the gravitational redshift effect by dark matter haloes show a good agreement with the predictions of General Relativity. \\end{itemize}"
    },
    "1208/1208.0587_arXiv.txt": {
        "abstract": "Radial trends of stellar populations in galaxies provide a valuable tool to understand the mechanisms of galaxy growth. In this paper, we present the first comprehensive analysis of optical--optical and optical--NIR colours, as a function of galaxy mass, out to the halo region ($8\\,\\rm R_e$) of early-type galaxies (ETGs). We select a sample of 674 massive ETGs ($M_\\star\\simgt 3\\times 10^{10}M_\\odot$) from the SDSS-based SPIDER survey.  By comparing with a large range of population synthesis models, we derive robust constraints on the radial trends in age and metallicity. Metallicity is unambiguously found to decrease outwards, with a measurable steepening of the slope in the outer regions ($\\rm R_e<R<8\\,R_e$). The gradients in stellar age are found to be more sensitive to the models used, but in general, the outer regions of ETGs feature older populations compared to the cores. This trend is strongest for the most massive galaxies in our sample ($M_\\star\\simgt 10^{11}M_\\odot$). Furthermore, when segregating with respect to large scale environment, the age gradient is more significant in ETGs residing in higher density regions. These results shed light on the processes leading from the formation of the central core to the growth of the stellar envelope of massive galaxies. The fact that the populations in the outer regions are older and more metal-poor than in the core suggests a process whereby the envelope of massive galaxies is made up of accreted small satellites (i.e. minor mergers) whose stars were born during the first stages of galaxy formation. ",
        "introduction": "The link between chemical and dynamical evolution in early-type galaxies (hereafter ETGs) has been studied for quite a long time through systematic measurements of colour gradients, which may lead us to understanding how the variations in the age and metallicity of the underlying stellar populations evolve as star formation proceeds through the history of the system \\citep{devac61,davies87,PVJ90,ferr05,LdC:09}.  The theoretical models suggest that metallicity and age gradients arise naturally from the processes leading to galaxy formation, like in the ``in-situ'' collapse of a proto-galactic gas cloud \\citep{carlberg84}.  The first rendition of this formation scenario, also known as monolithic collapse \\citep{els62,lar74,lar75}, predicted steep metallicity gradients, while, a revised monolithic model for the formation of ETGs led to shallower metallicity gradients \\citep[e.g.][]{pip08,pip10}. Simulations \\citep[see e.g.][]{koba04} indicate that galaxy mergers will weaken the metallicity gradient, and lead to significant levels of star formation in the galaxy centre. { More recently, cosmological simulations have brought to our attention the importance of cold accretion, as a viable way to form galaxies  ``in-situ'', resembling the monolithic formation scenario~\\citep[e.g.][]{keres05}.} It is still a matter of debate whether the metallicity gradient we measure today is a product of an old stellar population -- reflecting mainly the initial conditions when the system collapsed -- or results from more recent star formation episodes. Therefore, systematic measurements of colour gradients are of crucial importance to distinguish among different formation models of early-type systems. A proper constraint on the gradients of the properties of the underlying stellar populations will help understand the mass assembly process in ETGs.  In the past, spectroscopic indices were used to measure metallicity in ETGs, revealing the existence of radial gradients ranging from $-$0.1 to $-$0.3~dex per decade \\citep{cmc93,davies93,mehlert03}. The number of ETGs for which the radial dependence of age and metallicity extends beyond one effective radius is small for today standards, albeit growing. This is due to the difficulty in obtaining spectroscopic measurements at surface brightness below the background sky level. However, the availability of multi-waveband samples for a vast number of sources remedy this situation, allowing us to estimate colour gradients up to several effective radii at low to moderate redshifts \\citep{wu05,TalvanDokkum:11}, and open the possibility of studying radial gradients at z$>$1 \\citep{Gargiulo:11, GUo:11}.  An important caveat is that colours have to be modelled by a stellar population synthesis code. This approach has been used in the past for small samples of ETGs \\citep[e.g.][]{PVJ90,SMG00}, while in the recent years, the study of large galaxy samples has mostly focused on the optical regime \\citep[e.g.][]{RBH:10,Tor:10, GP11}, where one has to contend with the age-metallicity-extinction degeneracy \\citep{wo94}. \\citet{SE94}, and more recently~\\citet{Roediger:11}, have found evidence for significant age gradients in ellipticals, while independent works \\citep[e.g.][]{HI01,mehlert03}, found negligible age gradients, showing that the interpretation of colour gradients is a difficult task.  {  Another  aspect of  studying  colour  gradients  in ETGs  is  to understand  how  this  quantity  relates to  other  observables  and relationships so that we can  build up a consistent galaxy formation and evolution framework. For  instance, although} ETGs obey specific scaling relations  (e.g. Fundamental  Plane), they may  well represent different  families according  to the  way mass  was  assembled during their  history  and  the  characterization  of how  mergers  may  have established   the  mass   assembly  is   of   unquestionable  interest \\citep[e.g.][]{Greene:12}. As  a first approximation,  we would expect mergers  to  significantly  affect  the well-known  scaling  relations involving   central  velocity   dispersion   and  stellar   population properties, unless mass is  accreted at very large radii.  Theoretical and observational work has shown  that the presence of tidal debris in the outskirts  of ETGs may provide  an essential piece  of evidence to distinguish among different assembly scenarios \\citep[e.g.][]{Duc:11}. So far, most of the stellar population work was carried out within the central  regions   and  mainly  restricted  to   the  optical  regime. Extending  the analysis  to the  outer regions  with a  combination of optical and  near-infrared data is  a key factor to  better comprehend how galaxies form and assemble their baryons.  { In this paper we  combine optical (Sloan Digital Sky Survey--Data Release  6,  SDSS--DR6;~\\citealt{ade08})  and  near-infrared  (UKIRT Infrared    Deep   Sky    Survey--Data   Release    4,   UKIDSS-DR4; \\citealt{Law07}) } data to study colour gradients in a sample of 674 ETGs located in different environments.  The data were analysed with a dedicated pipeline  that uses the package  2DPHOT \\citep{LBdC08}.  The high  quality image  of both  optical and  NIR surveys  enabled  us to determine reliable  colour gradients out to eight  effective radii, as shown below. All systematic effects were carefully taken into account, including background subtraction; stacking  of the colour profiles and the effect of the tails of the  PSFs in the redder bands, all of which are crucial  for an accurate measurement  of the colours  in the outer regions.  The  results presented here confirm  previous findings about the metallicity gradients,  and show that age gradients  may not be as negligible as some papers have reported in the past.  The layout of the paper is as follows. In Sec.~\\ref{sec:samples} we present our sample of ETGs. Sec.~\\ref{sec:profiles} deals with the methodology followed to determine the colour profiles, and discusses possible sources of systematics. In Sec.~\\ref{sec:col_prof_mass} we present the median-stacked colour profiles of ETGs, from $g-r$ through $g-K$, out to a maximum galactocentric distance of $8\\,\\rm R_e$, for different galaxy mass bins. Sec.~\\ref{sec:fitting} describes the fitting of the observed profiles with theoretical models based on different, state-of-the-art, stellar population synthesis codes. In Sec.~\\ref{sec:SP_Age_Z} we show the results of the stellar population fits, i.e. how age and metallicity are found to vary from the centre to the external regions of ETGs.  We also discuss the contribution of internal reddening. Sec.~\\ref{sec:environment} shows how the inferred age and metallicity profiles are affected by galaxy environment, while in Sec.~\\ref{sec:summary} we summarize the main findings of this paper.  Throughout the paper, we adopt a standard $\\Lambda$CDM cosmology with $\\rm H_0 \\! = \\!  75\\, km \\, s^{-1} \\, Mpc^{-1}$, $\\Omega_{\\rm m} \\!  = \\!  0.3$, and $\\Omega_{\\Lambda} \\! = \\! 0.7$.   \\begin{figure*} \\begin{center} \\includegraphics[height=140mm]{f1.ps} \\end{center} \\caption{Comparison of the median surface brightness profiles of ETGs in different bins of stellar mass, $M_{\\star}$ (increasing from left to right), and axis ratio, $b/a$ (decreasing from top to bottom). For each panel (i.e. a given bin of $M_{\\star}$ and $b/a$), black curves plot the $r$-band surface brightness profiles of all galaxies in that bin. The profiles are shown as a function of the normalized galactocentric distance $\\rm R/R_e$, where $\\rm R_e$ is the $r$-band effective radius.  For each panel, the solid red curve is obtained by median-stacking all the single profiles, with dotted curves marking the $\\pm 1\\sigma$ scatter around the median profiles. For each bin of $M_{\\star}$, the median profile that corresponds to highest $b/a$ (top panel) is repeated in all panels (from top to bottom) as a coloured dashed curve (blue, magenta, and orange for left, middle, and right panels, respectively).  } \\label{fig:profs} \\end{figure*} ",
        "conclusions": "\\label{sec:summary}  We have used optical (SDSS-DR6) and near-infrared (UKIDSS-DR4) data for 674 massive ETGs to measure, for the first time, colour gradients out to $\\sim 8\\,\\rm R_{e}$, and study how ETGs assemble their mass. The sample comprises systems found in galaxy groups and in the field, with stellar masses ranging from $\\sim 3\\times 10^{10}\\rm M_{\\odot}$ to $\\sim 7\\times 10^{11}\\rm M_{\\odot}$. The critical systematic effects that may affect our conclusions have been carefully taken into account, and we show that none of them (e.g. sky subtraction, colour dependence of the PSF outer wings, and the approach adopted to derive the colour profiles) introduce any systematic biases. The stacked colour profiles show remarkable linearity out to $\\rm R \\sim 8\\,\\rm R_{e}$, getting bluer in the outer regions, irrespective of the mass bin, with the exception of $g - r$ where a curvature is apparent. Nevertheless, this effect is small and does not affect the stellar population analysis, as the deviations are comparable to the random errors at the radii where the curvature is detected. For the most massive galaxies, the stacked $g-r$ colour profile is compared with that recently obtained by~\\citet{TalvanDokkum:11}, finding good agreement.  Colour   profiles  are  translated   into  SSP-equivalent   $Age$  and metallicity gradients  using a  variety of stellar  population models, assuming either  a Chabrier  or Salpeter IMF,  as detailed  in Section 5.2.   In particular,  we exploit  the new,  state-of-the-art, stellar population models of Charlot  \\& Bruzual (2013, in preparation), based on different  stellar libraries  (IndoUS, Miles, STELIB,  BaSeL). { Errors  on  both age  and  metallicity  are  carefully estimated  by accounting  for uncertainties  on  observed colors.}   We find  very unambiguously  that metallicity  decreases from  $0.1$ out  to $8\\,\\rm R_{e}$ (negative gradient), regardless of the mass bin considered (see Tab.~6). This result is  independent of the adopted stellar population model and  IMF.  In  addition, we find  that the  metallicity gradient tends to steepen at  large galacto-centric radii.  The significance of this  result changes with  the adopted  stellar population  models and galaxy  mass bin (being  more significant  at low  mass).  As  for the $Age$ parameter, the situation is more complex.  From $0.1$ to $1\\,\\rm R_{e}$,  we find  that  for intermediate-  and  high-mass ETGs,  $Age$ increases  with  radius,  while  at  low mass,  no  clear  trend  with galacto-centric  distance is  detected, consistent  with  our previous results  (see e.g.~\\citealt{LF:11} and  references therein).   The new result of the  present work is that in the outer  regions, from $1$ to $8\\,\\rm R_{e}$, the stellar populations of ETGs are even older than in their centres.  This  is more significant for galaxies  in the highest stellar mass bin (i.e.  $M_\\star\\simgt 10^{11}M_\\odot$). We emphasize that using models with extended star formation histories, like exponentially declining ($\\tau$) and finite burst models, instead of SSPs, yields similar results within the quoted error bars. We also confirm that the presence of uniformly distributed internal reddening (dust) does not affect the trends of age and metallicity with galactocentric distance.  Finally, we analyze how the age and metallicity profiles depend on the environment  where  galaxies reside.   Group  ETGs  have positive  age gradients out to $\\sim 8\\,\\rm R_e$, with a very old stellar population in  the outskirts.   Their metallicity  gradients steepen  at  $\\rm R> 1-2\\,\\rm R_e$.  In  constrast, field ETGs do not  have significant age gradients in the inner region,  while at $\\rm R\\simgt 1\\,\\rm R_e$, the $Age$ shows  either a flat  trend or increases outwards  (depending on the  model).  The  metallicity  gradient  of field  ETGs  is found  to steepen  significantly in  the  outskirts only  for low-mass  systems. Notice that the results for group galaxies are essentially the same as those for  the whole sample of  ETGs (see above),  consistent with the fact  that most  of our  sample  ($\\sim 80\\%$)  comprises ETGs  either residing in (or close  to) galaxy groups (see Sec.~\\ref{sec:samples}). Our  findings  point  to  the  importance  of  aperture  effects  when comparing  the  stellar  population   content  of  ETGs  in  different environments, as in the galaxy outskirts, group galaxies are older (or coeval, depending on  the model) and less metal-rich  than field ETGs, while the opposite holds in the central regions, { where group ETGs are  found to  be more  metal-rich than  their  field counterparts}. This might explain, at least  in part, why different results have been reported in  the literature about the environmental  dependence of the metal  content  of  ETGs  (see  Paper  III),  with  different  studies reporting   field  ETGs   to   be  more   metal-rich~\\citep{Thomas:05, deLaRosa:07,    Kunt:02,   Clemens:09,   Zhu:10},    as   metal-rich as~\\citep{BERN:06,  Annibali:07}, or even  more metal-poor  than their cluster counterparts~\\citep{GALL:06}.  { We have  shown here that the stellar populations  in the halos of ETGs are consistently older and more metal poor than in their cores. This trend depends on mass  and, to a lower degree, environment, but overall, the  age and metallicity  radial gradients are  similar for all  ETGs   in  the  sample   (i.e.   more  massive   than  $3\\times 10^{10}$M$_\\odot$).  Within  the core of the  galaxies, the observed metallicity gradients  -- between \\nablazi$=-0.3$ and  $-0.4$ -- are consistent  with  the simulations  of  \\citet{koba04} for  non-major merger  systems.  For galaxies  where major-merging  is significant, her simulations give shallower metallicity gradients, around $-0.2$. It  is worth  noting  that our  metallicity  gradients get  slightly steeper with  increasing mass, supporting  the scenario of  fast and early formation  for the most massive ETGs.   Extending the analysis radially  out to  $8{\\rm R}_e$  allows us  to explore  the different channels of  star formation and  assembly in ETGs.  Our  results are consistent with the scenario  of \\citet{oser10}, where two phases of formation  operate:  an early-phase  of  in-situ  star formation  -- building  the core of  the galaxy  -- followed  by the  accretion of small  satellites.   The  recent simulations  of  \\citet{Lackner:12} quantify  in  more detail  the  buildup  of  the stellar  component, concluding that major mergers  do not dominate the accretion history of massive  galaxies, with smaller systems consisting  of older, and more  metal-poor stellar  populations  building up  the outer  halo, consistent  with  our   analysis.   Their  simulations  suggest  age (metallicity)  differences of  $\\Delta  t\\sim 2.5$~Gyr  ($\\Delta\\log Z\\sim -0.15$\\,dex) between accreted  stars and those formed in-situ. This  process is  fully consistent  with the  observed  evolution of massive  early-type  galaxies  on  the mass-size  plane  \\citep[see, e.g.][]{daddi05,truj06,truj11,vdk08},  where the  core  is quickly formed during  an early phase  \\citep[see, e.g.][]{ig:09b}, followed by   the  growth  of   the  halo   via  minor   mergers  \\citep[see, e.g.][]{naab09}.  }    \\appendix"
    },
    "1208/1208.5477_arXiv.txt": {
        "abstract": "We search the American Association of Variable Star Observers (AAVSO) archives of the two best studied dwarf novae in an attempt to find light curves for long outbursts that are extremely well-characterized. The systems are U Gem and SS Cyg. Our goal is to search for embedded precursors such as those  that have  been found recently in the high fidelity {\\it Kepler}  data for superoutbursts of some members of the SU UMa subclass of dwarf novae. For the vast majority of AAVSO data, the combination of low data cadence and large errors associated with individual measurements precludes one from making any strong statement about the shape of the long outbursts. However, for a small number of outbursts, extensive long term monitoring with digital photometry yields high fidelity light curves.  We report the discovery of embedded precursors in two of three candidate long outbursts. {\\it This is the first time that such embedded precursors have been found in dwarf novae above the period gap,} and reinforces van Paradijs' finding that long outbursts in dwarf novae above the period gap and superoutbursts in systems below the period gap constitute a unified class. The thermal-tidal instability to account for superoutbursts in the SU UMa stars predicts embedded precursors only for short orbital period dwarf novae, therefore the presence  of embedded precursors in long orbital period systems $-$  U Gem and SS Cyg $-$ argues for a more general mechanism to explain long outbursts. ",
        "introduction": "{\\it Kepler} observations of short orbital period dwarf novae (below the ``period gap'') have given us for the first time clear evidence for the presence of embedded precursors, or ``failed'' outbursts, at the beginning of  superoutbursts of SU UMa systems (Cannizzo et al. 2012). This behavior had been partially seen before in some systems with fragmentary data, but is now quite clear. Thanks to the high fidelity AAVSO data of the past $\\sim$10 yrs, we report the discovery of such precursors also in two systems {\\it above} the period gap. This may have consequences for theoretical models for the long outbursts and the superoutbursts.   Cataclysmic variables (CVs $-$ Warner 1995ab) are semi-detached interacting binaries containing  a Roche-lobe filling K or M secondary that transfers matter to a white dwarf (WD). CVs show a ``gap'' between $P_h\\sim$$2$ and 3 (where $P_h = P_{\\rm orbital}/1$ hr) during which time the secondary star loses contact with its Roche lobe and mass transfer ceases as the systems evolve to shorter orbital periods. Thus at $P_h\\simeq3$ the binary becomes fully detached. At $P_h\\simeq2$ the secondary  comes back into contact with its Roche lobe and mass transfer resumes. Systems can also be ``born'' in the gap, so it is not completely empty. For $P_h \\la 2$ angular momentum loss from the binary is thought to be due solely to gravitational radiation. The CV subclass of dwarf novae (DNe) also have semi-periodic outbursts. SU UMa stars are DNe below the period gap exhibiting  short, normal outbursts (NOs) and superoutbursts (SOs). SOs show superhumps which are modulations in the light curve at periods slightly exceeding the orbital  period; superhumps are the defining property of SOs.    DNe outbursts are thought to be due to a thermal limit cycle accretion disk instability (Smak 1984) in which material is accumulated in quiescence and then dumped onto the WD during outburst. During short outbursts in longer period DNe, a  few percent of the stored mass is accreted, and during long outbursts a significant fraction $\\sim$0.2 of the stored mass is accreted. For the SU UMa stars, a SO  is thought to accrete   $\\ga0.7-0.8$ of the stored mass. Although the accretion disk is never in steady state during the limit cycle, it is close to steady state during long outbursts.    U Gem and SS Cyg are the two DNe with the most complete AAVSO coverage. U Gem ($P_h=4.25$) was discovered by John Russell Hind in 1855 (Hind 1856), and SS Cyg ($P_h=6.60$) by Louisa D. Wells in 1896 (Wells 1896). To date, U Gem has $\\ga$115,000 observations, and SS Cyg has $\\ga$455,000. Figure 1   % shows the number of individual observations $n_{\\rm obs}$  for each 24 hr interval versus time for each long term light curve.       \\begin{figure} \\centering \\epsscale{1.0} \\includegraphics[scale=0.45]{figure1.jpg} \\vskip -0.05cm \\figcaption{ The number of observations entering into a daily mean $n_{\\rm obs}$ for the historical AAVSO light curves of U Gem ({\\it top panel}) and  SS Cyg  ({\\it bottom panel}). } \\smallskip \\end{figure}       Recent observations by {\\it Kepler} have revealed details of the outburst behavior that have been partially seen previously\\footnote{For example, Figure 3.36 from Warner (1995b) shows precursors of varying depths in SOs of VW Hyi, but the light curves are smoothed, filled versions based on fragmentary data.}, but which can now be studied in much  greater detail (Cannizzo et al. 2012). Of particular interest for this work is the presence of an embedded ``failed'' NO at the start of  a SO. Within the context of the thermal-tidal instability (TTI) model for SOs (Osaki 1989; Ichikawa \\& Osaki 1992; Osaki 2005, see his Figures 3 and 4), embedded precursors are understood as being due to a temporary squeezing of the outer accretion disk by increased effective tidal forces due to the onset of a tidal instability.  This  instability drives the outer disk between circular and eccentric shapes when the outer edge of the disk expands beyond the point of 3:1 resonance with the binary orbital period, i.e., the point at which $2\\pi/\\Omega$ around the WD equals $P_{\\rm orbital}/3$ (Whitehurst 1988). The discovery of the tidal instability by Whitehurst led Osaki to propose the TTI, which combines the accretion disk thermal limit cycle instability with the tidal instability. A necessary condition for the tidal instability is that the mass ratio $q\\equiv M_2/M_1 < 0.25$ so that for long outbursts, in which a substantial amount of matter is stored in the disk, the presence of high viscosity material can expand the outer disk radius beyond the 3:1 radius. Thus superhumps are seen only in short orbital period DNe.    An important point is that, if the TTI model only applies to low $q$ binaries, and therefore only to short orbital period DNe, one expects it not to apply to systems longward of the period gap (since for Roche-lobe filling main-sequence donors, systems with $q \\la 0.25$ have $P_h \\la 2.5$ hr). Therefore, the presence of embedded precursors in systems above the period gap, should they exist, would argue for a more general physical mechanism for long outbursts in all DNe. Previous studies of DNe outbursts using amateur data have proved useful in delineating timescales and constraining models (e.g., Campbell  \\& Shapley 1940, Sterne,  Campbell  \\& Shapley 1940, Bath \\& van Paradijs 1983, van Paradijs 1983, Szkody  \\& Mattei 1984, Cannizzo \\& Mattei 1992, 1998, Ak, et al. 2002, Simon 2004). Van Paradijs (1983) studied a sample of DNe spanning the period gap and found that short outburst durations increase with orbital period, whereas long outburst durations are relatively constant with orbital period. He proposed that the relation of SOs to NOs for DNe below the period gap and of long outbursts to NOs for DNe above the period gap are equivalent; superoutbursts are, in his view, just long outbursts seen in short orbital period DNe. This finding was amplified by Ak et al. (2002) using a larger sample.   For completeness we note that there have been prior critical analyses of the TTI model, from both theoretical and observational perspectives. Schreiber et al. (2004) compare numerical time dependent models of the TTI with an alternative model $-$ enhanced mass transfer (EMT) from the secondary star. Their arguments are somewhat general, however, and although they conclude in the end that the EMT is favored, they note ``we have not proven the EMT [model] to be correct nor the TTI [model] not to work.'' Also, in a series of papers Smak (2009abcd) presents  arguments against the standard interpretation of superhumps in the SU UMa stars as being due to a precessing, eccentric disk\\footnote{The most recent numerical simulations of an accretion disk subject to the 3:1 instability  reveal a disk shape that alternates between eccentric and circular (Montgomery 2012ab), therefore it may be more difficult than previously thought to detect an (instantaneously) eccentric disk shape in an eclipsing system.}. Thus from his perspective the entire physical basis for the TTI would be called into question. Smak also favors the EMT model. Cannizzo et al. (2012) study   the long term {\\it Kepler} light curves of two SU UMa stars  and argue  that the tendency of recurrence times for normal outbursts to exhibit sometimes a local maximum half way between   two superoutbursts argues against the TTI model, wherein one expects a monotonically increasing series of recurrence times  between  two superoutbursts (Ichikawa, Hirose, \\& Osaki 1993, see their Figure 1; Osaki 2005,  see  his Figure 3). ",
        "conclusions": "By examining the best AAVSO data for U Gem and SS Cyg we have found evidence for embedded failed outbursts at the start of two out of three long outbursts for which such an exercise is feasible. These data are from digital photometry. For the 2005 U Gem outburst, the possibility of a time variable calibration to explain the embedded precursor seems unlikely since the light curve is comprised of data from several observers. For the 2005 SS Cyg outburst, the precursor stands out more strongly. For other data in the long term light curves, the quality is not good enough  for one to be able to make any clear statement about the detailed outburst shape.   The thermal-tidal instability model was developed to account for the short orbital period SU UMa stars which have $q < 0.25$. Since CVs have Roche-lobe  filling secondary stars, this constrains the secondary star mass to be $\\sim$$0.1\\msun P_h$, unless the star has been driven considerably out of thermal equilibrium due to mass loss. Therefore DNe above the period gap such as U Gem and SS Cyg with orbital periods between 4 and 7 hr cannot possibly satisfy $q < 0.25$ for $M_{\\rm WD}\\simeq \\msun$, and therefore the thermal-tidal model should not be a physical ingredient of  their outbursts. Two of three long outbursts in DNe longward of the period gap which have the requisite AAVSO data fidelity show (i) the initial failed outburst embedded at the start, (ii) the slow decay consistent with viscous decay, and (iii) the faster decay consistent with thermal decay. These three characteristics are also manifested in the  long outbursts in DNe below the period gap $-$  i.e., the superoutbursts in the SU UMa stars. Therefore Occam's razor would seem to demand a common explanation which does not depend on $q$. The only difference between the two types of outbursts is the presence of superhumps in the superoutbursts, thus the association of  superhumps with  superoutbursts appears  to be associative rather than causal.   Finally, it is worth noting that U Gem exhibited an unusually long outburst in October 1985 with a duration of $\\sim$39 d (Cannizzo et al. 2002). The usual duration of  long outbursts in U Gem is $\\sim$12 d. It is unclear how the existence of this long outburst impacts our conclusions. In the context of the standard accretion disk limit cycle for DN outbursts (Smak 1984), the cooling front responsible for terminating the flat-topped portion of a long outburst must have been prevented from propagating, perhaps due to an unusually large amount of stored disk mass. Cannizzo et al. (2002) determined a viscous decay time $26\\pm6$ d mag$^{-1}$ from the AAVSO light curve. Given that virtually all other long outburst  decays in U Gem and SS Cyg do not last long enough for the decay slope to be measurable, this one long outburst  gives us our only measure of the viscous time scale in a DN longward of the period gap.    \\smallskip  We acknowledge the dedication and perseverance of the thousands of observers contributing data to the AAVSO International Database. We thank Allen Shafter for useful comments.     \\def\\mnras{MNRAS} \\def\\apj{ApJ} \\def\\apjs{ApJS} \\def\\apjl{ApJL} \\def\\aj{AJ} \\def\\araa{ARA\\&A} \\def\\aap{A\\&A} \\def\\aapl{A\\&AL} \\def\\pasj{PASJ}  \\vfil\\eject"
    },
    "1208/1208.0314_arXiv.txt": {
        "abstract": "MOSFIRE is a new multi-object near-infrared spectrometer for the Keck 1 telescope with a spectral resolving power of R$\\sim$3500 for a 0.7$^{\\prime\\prime}$ slit (2.9 pixels). The detector is a substrate-removed 2K x 2K HAWAII 2-RG HgCdTe array from Teledyne Imaging Sensors with a cut-off wavelength of 2.5 $\\mu$m and an operational temperature of 77K. Spectroscopy of faint objects sets the requirement for low dark current and low noise. MOSFIRE is also an infrared camera with a 6.9$^{\\prime}$ field of view projected onto the detector with 0.18$^{\\prime\\prime}$ pixel sampling. Broad-band imaging drives the requirement for 32-channel readout and MOSFIREs fast camera optics implies the need for a very flat detector. In this paper we report the final performance of the detector selected for MOSFIRE. The array is operated using the SIDECAR ASIC chip inside the MOSFIRE dewar and v2.3 of the HxRG software. Dark current plus instrument background is measured at $<$0.008 e$^{-}$ s$^{-1}$ pixel$^{-1}$ on average. Multiple Correlated Double Sampling (MCDS) and Up-The-Ramp (UTR) sampling are both available. A read noise of $<$5e$^{-}$ rms is achieved with MCDS 16 and the lowest noise of 3e$^{-}$ rms occurs for 64 samples. Charge persistence depends on exposure level and shows a large gradient across this detector. However, the decay time constant is always $\\sim$660 seconds. Linearity and stability are also discussed. ",
        "introduction": "\\label{sec:intro}  ~~~~ MOSFIRE is a near-IR spectrograph with four atmospheric bands: Y(0.97-1.12$\\mu$m), J(1.15-1.35$\\mu$m), H (1.46-1.81$\\mu$m), or K(1.95-2.39$\\mu$m).  It is located at the Cassegrain focus of the Keck-1 telescope on Mauna Kea, Hawaii.  The optical design provides both imaging and multiple-object spectroscopy over a field of view (FOV) of $6.14^\\prime$ x $6.14^\\prime$ with a resolving power ($\\frac{\\lambda}{\\Delta\\lambda}$) of R $\\simeq$ 3,500 for a slit width of $0.7^{\\prime\\prime}$ (2.9 pixels along the dispersion).  The detector is a sensitive 2K x 2K H2-RG HgCdTe array with a 2.5 $\\mu$m cut-off.  A unique feature of MOSFIRE is that its multiplex advantage of up to 46 slits is achieved using a cryogenic Configurable Slit Unit, or CSU, which has been developed in collaboration with the Swiss Center for Electronics and Micro Technology (CSEM).  The CSU is reconfigurable under remote control in less than 6 minutes without any thermal cycling of the instrument.  Slits are formed by moving opposing bars from both sides of the focal plane.  An individual slit has the length of $7.0^{\\prime\\prime}$, but bar positions can be aligned to make longer slits.  When the bars are moved apart to their full extent and the grating is changed to a mirror, MOSFIRE becomes a wide-field imager with a fine pixel scale of $0.18^{\\prime\\prime}$ per pixel.  MOSFIREs optical design is predicated on the 18 $\\mu$m pitch of the Teledyne H2-RG detector.  High quantum efficiency (QE $\\ge 65 \\%$) is required to achieve our instrument throughput goals.  The on-sky measured detection rates between OH lines in the darkest parts of the Y, J, H, and K bands are, respectively, 0.3, 0.3, 0.6, 0.4 $e^{-}$/s/pixel for a $0.7^{\\prime\\prime}$ slit assuming $\\sim$3 pixels sampling in the dispersion direction (R$\\simeq$ 3,500) for optimistic conditions.  A dark current of $<$0.03 $e^{-}$/s/pixel and an effective read noise below $\\sim5 e^{-}$ would result in background-limited performance for exposures longer than 90 seconds.  A typical spectroscopic integration time is 180 seconds in the K band based on the experience of the MOSFIRE team with the OH variability.  As stated in the MOSFIRE Design document, the limiting magnitudes (Vega) for a S/N ratio of 10 in 1000 seconds in spectroscopic mode at R$\\simeq$3500 are J = 20.4, H=20.1, and K = 18.6.  To achieve these limits we assume 0.05 $e^{-}$/s dark current, an effective read noise of 4 $e^{-}$/pixel, and the background appropriate for spectral regions between OH lines, evaluated over a 3 pixel resolution element and assuming a 0.5 arcsec$^{2}$ extraction aperture.  Imaging limits assume a 3 x 3 pixel (0.54$^{\\prime\\prime}$) aperture for a point source under good seeing conditions \\cite{spie2010}.   When it was selected for MOSFIRE in 2006, the H2-RG and especially the SIDECAR ASIC were recently developed technology from Teledyne Imaging Sensors. It was not known for certain at that time if the ASIC could achieve the kind of low noise performance typical of other controllers. To mitigate the risk and demonstrate that the detector and ASIC would meet the MOSFIRE sensitivity requirements, we undertook extensive testing of two engineering-grade devices and our science-grade detector. We worked closely with Teledyne on several iterations of ASIC hardware and software to achieve the final configuration.  The engineering-grade detectors were tested over a period of approximately two years.  The science-grade detector was then tested for an additional one and a half years before MOSFIRE achieved \\textquotedblleft first light\" at Keck.  Our results show that the MOSFIRE detector has very low dark current, low noise, excellent quantum efficiency across our wavelength range and good uniformity. It is also a very flat device, which is important for MOSFIREs fast camera beam. Reference pixels are effective in eliminating the readout shifts between the 32 outputs. The only behavior that may limit performance is charge persistence (residual images) from previous exposures. In the MOSFIRE detector there is an order of magnitude change in the charge persistence levels across the array. Our tests and analyses are described in this paper. ",
        "conclusions": "~~~~The H2-RG selected for MOSFIRE has proven to be an outstanding device in many respects. Measurements have shown low noise, low dark current, high QE and good uniformity. The chip is very flat and we have demonstrated that the point spread function is uniform across the field of view. Charge persistence may cause a problem, especially when changing from imaging to spectroscopic mode within a single night. Users will need to exercise care and allow sufficient decay time when changing modes. MOSFIRE will be available to the Keck community after August 2012 \\cite{spie2012}."
    },
    "1208/1208.3298_arXiv.txt": {
        "abstract": "We consider first order perturbation theory for a non-minimally coupled inflaton field without assuming an adiabatic equation of state. In general perturbations in non-minimally coupled theory may be non-adiabatic. However under the slow-roll assumptions the perturbation theory may look like adiabatic one. We show in the frame-work of perturbation theory, that our results of spectral index and bound on no-minimal coupling parameter agree with the results obtained using the adiabatic equation of state by the earlier authors. ",
        "introduction": "Inflationary paradigm has become extremely useful in solving many problems with the standard big-bang theory and very successful in predicting the fluctuations in the observed cosmic microwave background radiation\\cite{Guth:1980zm,Riotto:2002yw}. However, the nature of the inflaton potential remains to be uncertain and as a consequence there has been a variety of models  through which the inflationary paradigm can be implemented\\cite{Linde:2007fr} . One of the important class of such models that has been extensively investigated in recent time is that of a non-minimally coupled scalar fields\\cite{Futamase:1989cn,Komatsu:1999mt,Sakai:1998rg,Faraoni:1996rf,Makino:1991dc,Komatsu:1998ju,CervantesCota:1994zf,Gupta:2009kk,Setare:2008mb}. In most of the model of the inflation the mass of the scalar field is considered to be around $10^{13}$~GeV and the extremely small value for the strength of the quartic self-coupling $\\lambda \\sim 10^{-13}$\\cite{Bezrukov:2007ep}. This high value of the mass is considered to be an evidence for the physics beyond the standard model. In Refs.\\cite{Fakir:1990ab} it was shown that in case of chaotic inflation model instead of minimally coupled theory if non-minimal coupling is taken then amplitude of density perturbation will constrain the ratio $\\frac{\\lambda}{\\xi^{2}}$ rather than $\\lambda$ Therefore in minimally coupled theory one can remove the tight constraint on the self coupling parameter $\\lambda$ by choosing a higher value of $\\xi$. It is customary to choose the value of $\\xi\\sim10^{3}$. However, an extremely interesting possibility of having a non-minimally coupled standard model Higgs field as an inflaton has been pointed out in Ref.\\cite{Bezrukov:2007ep}. In this model for a sufficiently large strength of the non-minimal coupling $1\\ll\\sqrt{\\xi}\\lll 10^{17}$, it is possible to have $\\lambda \\approx 1$. Inert Higgs doublet has also been studied which gives the scale invariant density perturbation\\cite{Das:2011wm}. It is well known  \\cite{Dicke:1962ab} that the non-minimally coupled theory(Jordan frame) can be transformed into a minimally coupled theory(Einstein frame) by a set of transformations of the metric and the field. It should be noted that a consistent formalism of first order perturbation theory \\cite{Hwang:1990ac,Hwang:1990ab, PhysRevD.54.1460,Hwang:1998ab,Sugiyama:2010ab} and various bound on $\\xi$ \\cite{Kaiser:1995nv,Kaiser:1994vs,Fakir:1990ab} has been studied  by several earlier authors. Further, we would like to say in Refs.\\cite{Hwang:1990ac,Hwang:1990ab, PhysRevD.54.1460,Kaiser:1995nv,Hwang:1998ab} a `gauge ready' approach was developed to study the first order perturbation theory. In this approach the variable under a gauge condition which removes the gauge mode completely can be considered as the gauge invariant one. It has been shown recently in Ref.\\cite{Brown:2011eh} that all other gauges except Newtonian gauge is being tangled in the transformation between these two frames.    \\par In this work we are studying the first order perturbation in the the context of non-minimally coupled inflationary theory. We would like to emphasize that in order to derive equation of motion for first order curvature perturbation in the previous works an adiabatic equation of state was used. But in general the adiabatic equation of state for non-minimally coupled theory may not be possible \\cite{PhysRevD.54.1460,Carloni:2006gy}. Recently the non-adiabatic evolution of curvature perturbation is shown to exist in the case of non-minimally coupled multi-field inflation scenario \\cite{White:2012ab}. Therefore, we believe that it is important to relax the assumption about the equation of state for a single field inflaton theory. In this work we calculate the first order curvature perturbation and the spectral index without assuming any equation of state. Our results are consistent with the results obtained by earlier workers.      As mentioned earlier , for the case of chaotic inflation in Jordan frame that the density perturbation on Jordan frame are constrained by the ratio $\\frac{\\lambda}{\\xi^{2}}$ and this may allow for the values $\\xi\\gg1$. Further it is to be noted that in Ref. \\cite{Kaiser:1994vs} in order to produce the Harrison-Zel'dovich spectrum ($n_{\\mr}\\sim1$) for the `new-inflation' scenario it was necessary to put an upper bound on $\\xi$. However no such bound was required for the chaotic inflation scenario. I this work we show that in general there exist two branches, in one branch $\\dot\\vp$ (inflaton velocity) is negative and $\\dot\\vp$ is positive in the other branch depending on the nature of the potential. For those classes of potential where $\\dot\\vp>0$ one must have an upper bound on $\\xi$. But the other branch may allow a large value of $\\xi$ for some classes of the potentials provided the slow-roll parameter $\\epsilon_{V}$ is very small. It seems that for power-law kind of potentials only $\\lambda\\vp^{4}$ can allow large value of $\\xi$ which is consistent with the already known results.          The action in Jordan frame is given by \\begin{equation} S=\\int{d^{4}x\\sqrt{-g}\\left[f\\left(\\varphi\\right)R-\\frac{1}{2}\\varphi_{;\\mu}\\varphi^{;\\mu}+V\\left(\\varphi\\right)\\right]}, \\label{eq:actionj} \\end{equation} in case of non-minimal coupling $f\\lsp\\vp\\rsp=1+\\xi\\vp^{2}$, where $\\vp$ is the scalar field, $\\xi$ is a constant, $R$ is the Ricci scalar and $V(\\vp)$ is the potential. Here we have considered $M_{pl}=1$. The field equation in Jordan frame is given as \\bea \\label{eq:fe} \\nn f\\left(\\varphi\\right)&&\\left(R_{\\mu\\nu}-\\frac{1}{2}g_{\\mu\\nu}R\\right)=\\frac{1}{2}\\left(\\varphi_{;\\mu}\\varphi_{;\\nu}-\\frac{1}{2}g_{\\mu\\nu}\\varphi_{;\\alpha}\\varphi^{;\\alpha}\\right)+\\\\ &&\\frac{1}{2}g_{\\mu\\nu}V\\left(\\varphi\\right)+f\\left(\\varphi\\right)_{;\\mu;\\nu}-g_{\\mu\\nu}\\Box f\\left(\\varphi\\right) \\eea Jordan frame action (\\ref{eq:actionj}) can be transformed into the Einstein frame action \\be S=\\int{d^{4}x\\sqrt{-\\hat g}\\left[\\frac{1}{2}\\hat R-\\frac{1}{2}\\sigma_{;\\mu}\\sigma^{;\\mu}+\\hat V\\left(\\sigma\\right)\\right]}, \\label{eq:actione} \\ee by transforming the metric and the scalar field in the following way, \\be \\hat g_{\\mu\\nu}=2f\\lsp\\vp\\rsp g_{\\mu\\nu}, \\quad \\lsp\\frac{d\\sigma}{d\\vp}\\rsp^{2}\\equiv\\frac{1}{2f}\\lsp1+3\\frac{f^{2}_{,\\vp}}{f}\\rsp. \\label{eq:trans} \\ee where $\\hat\\quad$ represents the Einstein frame and $\\hat V(\\sigma)=\\frac{V(\\vp)}{4f^{2}}$. Using \\ref{eq:trans} the field equation in Eq.(\\ref{eq:fe}) can be shown to be transformed into the Einstein equations. This is true for the perturbed field equations also. In Einstein frame the equation of motion for the comoving curvature perturbation $\\hat{\\mr}=\\hat\\psi+\\hat{\\mh}\\frac{\\delta\\sigma}{\\sigma^{\\prime}}$ can be written as \\be \\hat\\mr^{\\prime\\prime}+\\lsb\\ln\\hat\\gamma\\rsb^{\\prime}\\hat\\mr^{\\prime}+k^{2}\\hat\\mr=0 \\label{eq:eomre} \\ee where $\\hat\\gamma=\\frac{\\hat a^{2}\\sigma^{\\prime 2}}{\\hat{\\mh}^{2}}$, $\\hat{\\psi}$ and $\\hat{\\mh}$ are metric perturbations and Hubble parameter in Einstein frame. In Einstein frame spectral index $\\hat n_{\\mr}$ can be expressed in terms of slow roll parameters as \\be \\hat n_{\\mr}-1=-4\\hat\\epsilon_{V}-2\\hat\\delta. \\label{eq:sie} \\ee where $\\hat\\epsilon_{V}=\\frac{3}{2}\\frac{\\hat{\\sigma}^{\\prime 2}}{\\hat a^{2}\\hat V}$ and $\\hat\\delta=1-\\frac{\\hat\\sigma^{\\prime\\prime}}{\\hat\\mh\\hat\\sigma^{\\prime}}$ are the slow roll parameters in the Einstein frame. ",
        "conclusions": ""
    },
    "1208/1208.0179_arXiv.txt": {
        "abstract": "{ We consider $1+3$  dimensional maximally symmetric Minkowski brane embedded in a $1+4$ dimensional maximally symmetric Minkowski background. The resulting $1+3$ dimensional effective field theory is of DBI (Dirac-Born-Infeld) Galileon type. We use this model to study the late time acceleration of the universe. We study the deviation of the model from the concordance $\\Lambda$CDM behaviour. Finally we put constraints on the model parameters using various observational data.} \\date{\\today} ",
        "introduction": "Providing a satisfactory explaination of the late time acceleration of the universe is one of the most challenging tasks for cosmologists and particle physicists at present \\cite{review}. The most studied approach invloves adding an exotic form of energy with negative pressure (termed as dark energy) in the energy budget of the universe. The cosmological constant  ( with $w= \\frac{p}{\\rho}=-1$) is the simplest candidate for such a component. Although the concordance $\\Lambda$CDM model is allowed by all current cosmological observations, it is also plagued by serious issues like fine-tuning and cosmic coincidence problems. Scalar field model \\cite{scalar} is another example of dark energy where the equation of state $w$ evolves with time which in turn helps to solve the cosmic coincidence by tracker- type evolution. This also restricts the form of the potential for the scalar fields.  Another interesting approach to explain the late time acceleration of the universe is to modify the gravity at large scale ( infra-red modification of gravity). DGP brane-world model \\cite{dgp} is one such example where the gravity is altered at large scales due to the slow leakage of gravitons from our observable universe ( modelled as an three brane) in to the higher dimensional bulk. The resulting Hubble equation can lead to late time acceleration in the observable universe. DGP model provides an interesting setup to modify gravity where one can construct a new four dimensional effective field theory  which contains nontrivial symmetry properties. These symmetries are due to the combination of five dimensional Poincare invariance and the brane paramatrization invariance. It has been shown that in this effective theory, there is a single scalar field $\\pi$ which represents the position of the three-brane ( our observable universe) in the higher dimensional bulk. The effective action contains a cubic self-interaction term of the kind $(\\partial \\pi)^2 \\Box\\pi$ together with normal canonical kinetic energy term. This cubic self interaction term has the property that it leads to second order equation of motion of $\\pi$. This term is also invariant under the {\\it Galilean Transformation}: \\begin{eqnarray} \\pi &\\rightarrow& \\pi +a\\nonumber\\\\ \\partial_{\\mu}\\pi &\\rightarrow& \\partial_{\\mu}\\pi + b_{\\mu} \\label{piequation} \\end{eqnarray}  \\noindent where $a$ and $b_{\\mu}$ are constants.  Due to the form of equation (\\ref{piequation}) $\\pi$ is often referred as Galileon field. Later, a four dimensional theory for the field $\\pi$ with Galileon symmetry was proposed that contains five terms which also had the inersting property that depsite the presence of higher derivatives in the action, the equations of motions are second order \\cite{lutty, covariant} .  This causes the theory to be free of any ghosts. Cosmology with such an action has been widely studied in recent times \\cite{galileon}.  On the other hand, Dirac-Born-Infeld (DBI) action contains the lowest order dynamics of a brane embedded in a higher dimensional spacetime. This gives an interesting setup to study inflation \\cite{dbiinf} and late time acceleration \\cite{dbiquint,dbiessence} of the universe. If the universe is indeed described in a brane world scenario, then they should share the symmetries in the DBI action. Galileon terms can be thought of as a subset of all the higher derivative terms usually expected to be present in any effective field theory of the brane and they will be suppressed by some cut of scale (of different powers). In a recent work, Rham and Tolley \\cite{rham} have constructed a general class of effective field theory which reproduces Galileon as well as the familiar DBI action under different limits. This has been extended by Goon et al. \\cite{goon} where they have  constructed a general class of effective field theory assuming that a 3-brane is moving in the higher dimensional bulk. Galileon theory arises as a special case of this setup.  Similar generalization of Galileon field to maximally symmetric spacetime using a de-Sitter slicing has also been done in \\cite{burrage}.  In this work, we study the late time acceleration of the universe in such a setup. We assume the simplest case where a maximally symmetric brane ( Minkowski Brane) is embedded in a maximally symmetric bulk (Minkowski bulk). We consider those three terms in the action which under small field limit generate the standard Galileon terms that arise in the DGP model under decoupling limit. We compare the cosmological behaviour in this setup to that of the concordance $\\Lambda$CDM universe. We also constrain our model parameters using currently available observational data.  We start by giving a general introduction on DBI-Galileon model in section 2. In section 3, we discussed the late time cosmology in this model. We explore the constraints on the model parameters using observational data in section 4. Section 5 deals the conclusions. ",
        "conclusions": "\\vspace{5mm}  In this work, we investigate the late time evolution of the universe in a DBI-Galileon Model where a Minkowski Brane is embedded in a Minkowski bulk. To simplify our analysis, we only keep the first three terms in the total action which under weak-field limit reproduce the standard Galileon terms present in the decoupling limit of DGP model.  We also keep a general potential term $V(\\pi)$ in the action instead of standard linear term and study different choices for $V(\\pi)$.  We assume the field $\\pi$ to be initially frozen due to large Hubble damping and that it behaves like a cosmological constant. However, with time the field slowly thaws out from the frozen state and starts deviating from $w=-1$. The deviation depends on the initial value of te slope of the potential, $\\lambda_{i}$. For smaller values of $\\lambda_{i}$, the evolution for all the potentials remain close to $\\Lambda$CDM throughout. As one increases the value of $\\lambda_{i}$, the evolution start deviating from $\\Lambda$CDM. We study the degeneracies for the different potentials using the statefinder hierarchies. Finally we constrain our model parameters using the recent observational data. We show that larger contribution from ${\\cal L}_{2}$ part results in larger deviation from $\\Lambda$CDM behaviour whereas larger contribution contribution from ${\\cal L}_{3}$ part restricts the models to behave more close to $\\Lambda$CDM. This is true for all the potentials.  Although we have not studied the complete action for the DBI-Galileon model, still this study gives some interesting observational consequences for first three terms of the full action. It will be worthwhile to study the observational consequences with full DBI-Galileon action and this will be our future aim."
    },
    "1208/1208.5330_arXiv.txt": {
        "abstract": "Conformal gravity theory can explain observed flat rotation curves of galaxies without invoking hypothetical dark matter. Within this theory, we obtain a generic formula for the sizes of galaxies exploiting the stability criterion of circular orbits. It is found that different galaxies have different finite sizes uniquely caused by the assumed quadratic potential of cosmological origin. Observations on where circular orbits might actually terminate could thus be very instructive in relation to the galactic sizes predicted here. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.2510_arXiv.txt": {
        "abstract": "The static electrical conductivity of non-ideal, dense, partially ionized helium plasma was calculated over a wide range of plasma parameters:  temperatures $1\\cdot 10^{4}\\textrm{K} \\lesssim T \\lesssim 1\\cdot 10^{5}\\textrm{K}$ and mass density $1 \\times 10^{-6} \\textrm{g}/\\textrm{cm}^{3} \\lesssim \\rho \\lesssim 2 \\textrm{g}/\\textrm{cm}^{3}$. Calculations of  electrical conductivity of plasma for the considered range of plasma parameters are of interest for DB white dwarf atmospheres with effective temperatures $1\\cdot 10^{4}\\textrm{K} \\lesssim T_{eff} \\lesssim 3\\cdot 10^{4}\\textrm{K}$.  Electrical conductivity of plasma was calculated by using the modified random phase approximation and semiclassical method, adapted for the case of dense, partially ionized plasma. The results were compared with the unique existing experimental data, including the results related to the region of dense plasmas. In spite of low accuracy of the experimental data, the existing agreement with them indicates that results obtained in this paper are correct. ",
        "introduction": "DB white dwarf atmospheres belong to the class of astrophysical objects which have been investigated for a long time and from various aspects \\citep{bue70,koe80,sta93}. In the previous period, the contribution of the authors of this paper was related to the research of optical properties of DB white dwarf atmospheres within the range of average effective temperatures $1\\cdot 10^{4}\\textrm{K} \\lesssim T_{eff} \\lesssim 2\\cdot 10^{4}\\textrm{K}$. So, the papers of \\citet{mih92b} and of \\citet{mih94,mih95} were dedicated to the research of continual absorption in the optical part of EM spectra, the paper of \\citet{mih03b} - to the research of chemi-ionization/recombination processes and of \\citet{ign09} - to the investigation of continual absorption in the VUV region of EM spectra.  Recently, the transport properties of helium plasmas, characteristic of some DB white dwarf atmospheres, attracted the authors' attention, first of all the electrical conductivity. Namely, the data on electrical conductivity of plasma of stars with a magnetic field or moving in the magnetic field of the other component in a binary system (see e.g. \\citet{zha09}, \\citet{pot09} and \\citet{rod09}) could be of significant interest, since they are useful for the study of thermal evolution of such objects (cooling, nuclear burning of accreted matter) and the investigation of their magnetic fields. For example, \\citet{kop70} and \\citet{kop73} studied electrical conductivity for stars of various spectral types, in order to investigate the magnetohydrodynamic differences in their atmospheres. Recently, \\citet{maz07} investigated He conductivity in cool white dwarf atmospheres, since the possibility of using these stars for dating stellar populations has generated a renewed interest in modelling their cooling rate \\citep{fon01}. Also, the transport processes occurring in the cores of white dwarfs (see e.g. \\citet{bai95} and numerous references therein) have been considered. Moreover, electrical conductivity was particularly investigated for solar plasma, since it is of interest for consideration of various processes in the observed atmospheric layers, like the relation between magnetic field and convection, the question of magnetic field dissipation and the energy released by such processes (see e. g. \\citet{kop70} and references therein). For example \\citet{fel93} investigated the role of electrical conductivity in the construction of a theoretical model of the upper Solar atmosphere, and \\citet{kaz06} considered the electrical conductivity of solar plasma for magnetohydrodynamic simulations of the solar chromospheric dynamo. Given that electrical conductivity plays an analogous role in other stars as well, it is of interest to investigate its significance, to adapt the methods for research into stellar plasma conditions and to provide the needed data.  An additional interest for data on electrical conductivity in white dwarf atmospheres may be stimulated by the search for extra-solar planets. Namely \\citet{jia98} have shown that a planetary core in orbit around a white dwarf may reveal its presence through its interaction with the magnetosphere of the white dwarf. Such an interaction will generate electrical currents that will directly heat the atmosphere near its magnetic poles. \\citet{jia98} emphasize that this heating may be detected within the optical wavelength range as H$_\\alpha$ emission. For investigation and modelling of mentioned electrical currents, the data on electrical conductivity in white dwarf atmospheres will be useful.  One of the most frequently used  approximations for consideration of transport properties of different plasmas is the approximation of \"fully ionized plasma\" \\citep{spi62,rad76,ada80,kur84,rop89,dju91,nur97,zai00,ess03}. It was shown that the electrical conductivity of fully ionized plasmas can be successfully calculated using the modified random-phase approximation (RPA) \\citep{dju91,ada94a,ada94b} in the region of strong and moderate non-ideality, while the weakly non-ideal plasmas were successfully treated within the semiclassical approximation (SC) \\citep{mih93,vit01}. In practice, even the plasmas with a significant neutral component are treated as fully ionized in order to simplify the considered problems, \\citep{rop89,ess98,zai00,ess03}. However, our preliminary estimates have shown that such an approach is not applicable for the helium plasmas of DB white dwarf atmospheres described in \\citep{koe80}, where the influence of neutral component cannot be neglected.  Therefore, an adequate method for calculations of electrical conductivity of dense, partially ionized helium plasmas is developed in this paper. This method represents a generalization of methods developed in \\citet{dju91} and \\citet{mih93}, namely modified RPA and SC methods, and gives a possibility to estimate the real contribution of the neutral component to the static electrical conductivity of the considered helium plasmas within a wide range of mass densities ($\\rho$) and temperatures ($T$).  The calculations were performed for helium plasma in the state of local thermodynamical equilibrium with given $\\rho$ and $T$ for $1\\cdot 10^{4}\\textrm{K} \\lesssim T \\lesssim 1\\cdot 10^{5}\\textrm{K}$ and $1 \\times 10^{-6} \\textrm{g}/\\textrm{cm}^{3} \\lesssim \\rho \\lesssim 2 \\textrm{g}/\\textrm{cm}^{3}$. The obtained results are compared with the corresponding experimental data (\\citet{min80,ter02,shi03}). For the calculations of plasma characteristics of DB white dwarf atmospheres the data from \\citep{koe80} were used. ",
        "conclusions": "In order to apply our results to the study of DB white dwarf atmosphere plasma properties, helium plasmas with electron ($N_{e}$) and atom ($N_{a}$) densities and  temperatures ($T$), characteristic of atmosphere models presented in the literature \\citep{koe80}, are considered here. So, the behaviour of $\\rho$ and $T$ for models with the logarithm of surface gravity $\\log g=$8 and effective temperature $T_{eff}=$12000K, 20000K and 30000K is shown in Fig.~\\ref{fig:DBmodels} as a function of Rosseland opacity $\\tau$. As one can see, these atmospheres contain layers of dense helium plasma. In order to cover the considered plasma parameter range reliably, we tested our method for calculation of the plasma electrical conductivity within a wider range of mass density $1 \\times 10^{-6} \\textrm{g}/\\textrm{cm}^{3} \\lesssim \\rho \\lesssim 2 \\textrm{g}/\\textrm{cm}^{3}$ and temperature $1\\cdot 10^{4}\\textrm{K} \\lesssim T \\lesssim 1\\cdot 10^{5}\\textrm{K}$.  The influence of neutral atoms on the electrical conductivity of helium plasma is shown in Fig.~\\ref{fig:withatoms}. In this figure the electrical conductivities for $T=15000$, $20000$ and $25000$K are given as functions of mass density $\\rho$. The range between the two vertical dashed lines corresponds to the conditions in the considered DB white dwarf atmospheres.  Two groups of curves, calculated using the expression (\\ref{eq:sigma0}), are presented in this figure: a) the dashed ones, obtained by neglecting the influence of atoms, i.e. with $\\nu_{ea}=0$; b) the full-line curves calculated with the influence of atoms included, i.e. with $\\nu_{ea}$ given by Eq.~(\\ref{eq:nu_ea}). First, one should note that the behaviour of these two groups of curves is qualitatively different: the first one increases constantly with the increase of $\\rho$, while the other group of curves decreases, reaches a minimum, and then starts to increase with the increase of $\\rho$. One could explain such behaviour of the electrical conductivity by the pressure ionization. This figure also clearly shows when the considered plasma can be treated as \"fully ionized\".  In Fig.~\\ref{fig:HeCond} we compare our values of the helium plasma conductivity, shown by full curves for $T=15000$ K, 20000 K and 25000 K within the region $5\\cdot 10^{-4}$g/cm$^{3}< \\rho < 2$g/cm$^{3}$, with the existing experimental data. Let us note that these experimental results are uniquely available for comparison so that, in spite of their low accuracy, the agreement with them gives the only possible indication that our results are correct.  Within the region $\\rho < 0.65$ g/cm$^{3}$, i.e. to the left of the vertical line in Fig.~\\ref{fig:HeCond}, there are experimental results from \\citet{ter99} ($\\circ$) and \\citet{shi03} ($\\triangledown$) where the temperature was determined with an error of less than $20\\%$, which are related to the temperature range 20000K-25000K. For $\\rho > 0.65$ g/cm$^{3}$, i.e. right of the vertical line in Fig.~\\ref{fig:HeCond}, are shown several values of the plasma conductivity, obtained by \\citet{ter02} for the temperature range 15000K-25000K. These experimental values are obtained with an experimental error $\\sim 50\\%$ and can be treated only as characteristic of this temperature region as a whole. These data at least indicate that our results for $\\rho > 0.65$ g/cm$^{3}$ lie in the correct domain of the electrical conductivity values.  The developed method was then applied to calculation of plasma electrical conductivity for the models of DB white dwarf atmospheres presented in Fig.~\\ref{fig:DBmodels}. The results of the calculations are shown in Fig.~\\ref{fig:DBcond}. First, let us note a regular behaviour of the static electrical conductivity which one should expect considering the characteristics of DB white dwarf atmospheres. Further, the electrical conductivity profiles presented in this figure show that, for the considered DB white dwarf models, plasma electrical conductivity changes over the domain of values where our results agree with the experimental ones (see Fig.~\\ref{fig:HeCond}). This indicates that the theoretical apparatus presented here may be adequate to be used for investigation of DB white dwarfs in the magnetic field of their partners in binary systems and magnetic white dwarfs.  \\begin{table} \\caption{Static electrical conductivity of helium plasma $\\sigma_{0}[1/(\\Omega m)]$} \\begin{tabular}{@{} r c c c c c c c c c @{} } \\hline\\hline \\multicolumn{1}{c}{} & \\multicolumn{9}{c}{$\\rho [g/cm^{3}]$} \\\\ \\cline{2-10} $T[K]$&5.00E-07 &1.00E-06 &5.00E-06 &1.00E-05 &5.00E-05 &1.00E-04 &5.00E-04 &1.00E-03 & 5.00E-03 \\\\ \\hline\\hline 8000 & 6.53E+00& 4.69E+00& 2.15E+00& 1.50E+00& 6.90E-01& 4.89E-01& 2.20E-01& 1.53E-01& 6.88E-02 \\\\ 9000 & 4.34E+01& 3.16E+01& 1.47E+01& 1.13E+01& 5.12E+00& 3.60E+00& 1.58E+00& 1.19E+00& 5.20E-01 \\\\ 10000 & 1.69E+02& 1.30E+02& 6.74E+01& 5.02E+01& 2.38E+01& 1.72E+01& 8.12E+00& 5.73E+00& 2.56E+00 \\\\ 12000 & 9.36E+02& 8.04E+02& 5.09E+02& 4.05E+02& 2.19E+02& 1.66E+02& 8.35E+01& 6.08E+01& 2.92E+01 \\\\ 14000 & 2.27E+03& 2.08E+03& 1.62E+03& 1.38E+03& 8.87E+02& 6.97E+02& 3.99E+02& 3.00E+02& 1.28E+02 \\\\ 16000 & 3.64E+03& 3.55E+03& 3.15E+03& 2.87E+03& 1.45E+03& 1.18E+03& 6.85E+02& 5.28E+02& 2.75E+02 \\\\ 18000 & 4.90E+03& 4.97E+03& 4.76E+03& 3.79E+03& 2.87E+03& 2.46E+03& 1.58E+03& 1.26E+03& 7.07E+02 \\\\ 20000 & 6.08E+03& 6.22E+03& 6.32E+03& 5.61E+03& 4.99E+03& 4.46E+03& 3.21E+03& 2.53E+03& 1.58E+03 \\\\ 25000 &         &         & 9.95E+03& 1.02E+04& 9.90E+03& 9.90E+03& 8.62E+03& 7.74E+03& 5.95E+03 \\\\ 30000 &         &         & 1.33E+04& 1.40E+04& 1.56E+04& 1.60E+04& 1.57E+04& 1.50E+04& 1.21E+04 \\\\ 35000 &         &         & 1.65E+04& 1.74E+04& 2.02E+04& 2.15E+04& 2.34E+04& 2.35E+04& 2.15E+04 \\\\ 40000 &         &         & 1.97E+04& 2.09E+04& 2.43E+04& 2.61E+04& 3.06E+04& 3.18E+04& 3.22E+04 \\\\ 45000 &         &         & 2.31E+04& 2.45E+04& 2.84E+04& 3.06E+04& 3.69E+04& 3.93E+04& 4.36E+04 \\\\ 55000 &         &         & 3.04E+04& 3.23E+04& 3.70E+04& 3.98E+04& 4.81E+04& 5.26E+04& 6.40E+04 \\\\ 65000 &         &         & 3.81E+04& 4.08E+04& 4.65E+04& 4.98E+04& 6.00E+04& 6.56E+04& 8.20E+04 \\\\ 75000 &         &         & 4.58E+04& 4.93E+04& 5.66E+04& 6.06E+04& 7.26E+04& 7.92E+04& 9.84E+04 \\\\ \\hline\\hline \\end{tabular} \\label{tab:ec} \\end{table}  In order to provide possibility for direct applications of our results to different theoretical investigations the values of static electrical conductivity $\\sigma_{0}$ of helium plasma in a wide range of $\\rho$ and $T$ are given in Table~\\ref{tab:ec}. This table covers plasma conditions for all models of DB white dwarf atmospheres presented in \\citet{koe80} and the corresponding values of $\\sigma_{0}$ were determined for the helium plasmas in the state of local thermodynamical equilibrium with given $\\rho$ and $T$.  The method developed in this paper represents also a powerful tool for research into white dwarfs with different atmospheric compositions (DA, DC etc.), and for investigation of some other stars (M-type red dwarfs, Sun etc.). Finally, the presented method provides a basis for the development of methods to describe other transport characteristics which are important for the study of all mentioned astrophysical objects, such as the electronic thermo-conductivity in the star atmosphere layers with large electron density, electrical conductivity in the presence of strong magnetic fields and dynamic (high frequency) electrical conductivity."
    },
    "1208/1208.4100_arXiv.txt": {
        "abstract": "Recent claims of a line in the Fermi-LAT photon spectrum at 130 GeV are suggestive of dark matter annihilation in the galactic center and other dark matter-dominated regions. If the Fermi feature is indeed due to dark matter annihilation, the best-fit line cross-section, together with the lack of any corresponding excess in continuum photons, poses an interesting puzzle for models of thermal dark matter: the line cross-section is too large to be generated radiatively from open Standard Model annihilation modes, and too small to provide efficient dark matter annihilation in the early universe.  We discuss two mechanisms to solve this puzzle and illustrate each with a simple reference model in which the dominant dark matter annihilation channel is photonic final states. The first mechanism we employ is resonant annihilation, which enhances the annihilation cross-section during freezeout and allows for a sufficiently large present-day annihilation cross section. Second, we consider cascade annihilation, with a hierarchy between $p$-wave and $s$-wave processes.  Both mechanisms require mass near-degeneracies and predict states with masses closely related to the dark matter mass; resonant freezeout in addition requires new charged particles at the TeV scale. ",
        "introduction": "\\label{sec:intro}  The existence of dark matter (DM) is one of the strongest pieces of evidence for physics beyond the Standard Model.  Searches in cosmic rays for evidence of DM annihilation or decays are a cornerstone of the experimental effort to detect DM. Monochromatic photon lines, though in most models a subdominant signal, provide a particularly clean astrophysical signal~\\cite{Bergstrom:1988fp, Bern:1997ng, Bergstrom:2004nr, Bertone:2009cb, Bertone:2010fn}.  Several recent analyses have claimed evidence for a distinct spectral feature in the Fermi-Large Area Telescope (LAT) \\cite{Atwood:2009ez} photon spectrum at around 130 GeV~\\cite{Bringmann:2012vr, Weniger:2012tx, Tempel:2012ey, Su:2012ft}, in regions near the galactic center. Evidence for this feature has also been reported in galactic clusters \\cite{Hektor:2012kc} and in non-associated sources \\cite{Su:2012zg}, although the latter claim remains contentious \\cite{Hooper:2012qc, Mirabal:2012za, Hektor:2012jc}. While the statistics are limited, the morphology of the signal may favor an explanation as annihilating DM \\cite{Cotta:2011pm, Buchmuller:2012rc, Hektor:2012kc}. The presence of an additional photon line at approximately 111 GeV \\cite{Su:2012ft, Su:2012zg} is highly suggestive, if also statistically limited, and would lend more credence to a particle physics explanation \\cite{Rajaraman:2012db}. The Fermi collaboration's own search for photon lines uses slightly different search regions and methodology and sets an upper limit marginally in conflict with the claimed signal~\\cite{Ackermann:2012qk}.  It remains to be established whether the excess is instrumental, astrophysical, or representative of an overly optimistic characterization of the systematic uncertainties in the galactic background \\cite{Boyarsky:2012ca, Aharonian:2012cs}.  However, standard thermal WIMPs are not capable of explaining the Fermi signal, and it is of interest to work out the necessary structure in DM models which could give rise to the 130 GeV line.  Any dark matter model for the Fermi 130 GeV line must account for two interesting facts. First, there is no evidence for an excess in the continuum photon spectrum, which strongly constrains DM annihilation into the usual SM annihilation channels ($f\\bar f$, $W^+W^-$, $ZZ$)~\\cite{Buchmuller:2012rc, Cohen:2012me, Cholis:2012fb, Huang:2012yf} or indeed into any charged final states.  Since in most thermal models DM-photon couplings are generated radiatively from the DM couplings to charged final states~\\cite{Cline:2012nw, Das:2012ys}, the typical line cross section is generically related to the cross-section for annihilation into charged modes $X,X^\\dag$ by \\beq \\label{eq:continuum} \\langle \\sigma v\\rangle_{\\gamma\\gamma} \\sim \\left(\\frac{\\alpha}{\\pi}\\right)^2 \\langle \\sigma  v\\rangle_{X X^\\dagger}\\,. \\eeq As the fragmentation and decay of the final states $X,X^\\dag$ give rise to a continuum photon spectrum $d\\Phi_\\gamma (E)/dA \\propto n_{DM}^2\\langle \\sigma v\\rangle_{X X^\\dagger} dN_\\gamma/dE $, if annihilation to charged states is open, the expected line flux is smaller than the continuum flux by several orders of magnitude. Models for the Fermi 130 GeV line must therefore explain the absence of annihilation into charged (or hadronic) modes.  This brings us to the second interesting fact.  The best fit cross-sections for the 130 GeV feature \\cite{Bringmann:2012vr,Weniger:2012tx} are more than an order of magnitude smaller than the expectation for ($s$-wave) thermal freezeout.  If annihilation to SM or charged modes is absent or suppressed as the continuum limits suggest, reconciling this inefficient annihilation to photons with the WMAP relic density $\\Omega_\\chi h^2 = 0.1109\\pm 0.0056$~\\cite{Larson:2010gs} requires either: (1) a nonthermal relic abundance~\\cite{Acharya:2009zt, Tulin:2012uq}; (2) non-photonic annihilation modes which are suppressed, possibly only post-freezeout, relative to the naive radiative scaling of Eq.~(\\ref{eq:continuum})~\\cite{Weiner:2012cb, Tulin:2012uq}; or (3) a mechanism to enhance the thermal annihilation cross-section to photonic final states in the early universe~\\cite{Griest:1990kh}.  In the present work we will study two mechanisms which give enhanced DM annihilation to photons while also obtaining the correct thermal relic abundance, and build simple reference models for both.  Our first example is resonant freezeout, where the presence of a resonance in the DM sector spectrum enhances the annihilation cross-section into photons during freezeout. Our second example introduces an intermediate annihilation mode, so that DM annihilation proceeds through a cascade decay of a non-photonic intermediate state, $\\bar{\\chi}\\chi\\to \\phi \\phi' \\to 4 \\gamma$. In this example, the relation between the cross-section necessary for thermal freezeout and the present-day gamma line cross section is explained by the interplay of $s$-wave and $p$-wave contributions to the annihilation.  In section~\\ref{sec:effective}, we perform an effective operator analysis of DM-photon couplings and demonstrate the need for new particles at mass scales comparable to the DM mass.  In section~\\ref{sec:resonance} we perform a detailed examination of resonant freezeout, and consider cascade decays in section~\\ref{sec:cascade}. Section~\\ref{sec:conclusion} contains our conclusions. ",
        "conclusions": "\\label{sec:conclusion}  We have explored the possibilities for thermal models for the Fermi-LAT 130 GeV gamma line excess where the dominant DM annihilation channel is into photons.  We consider two mechanisms, (1) models where a resonance in the DM spectrum enhances the annihilation rate during freezeout and to a lesser extent at the present day, and (2) models where DM annihilates to a new intermediate state which subsequently decays to photon pairs.  Here the interplay of the $s$-wave and $p$-wave annihilation is responsible for reconciling the necessary cross-section at freezeout with the observed cross-section today. For the resonance model, charged fermions at the TeV scale are predicted and are accessible at the Large Hadron Collider, which could be interesting if the excess in the Higgs diphoton decay channel persists~\\cite{AtlasHiggs,CMSHiggs}.  Both of the classes of models considered in this paper require coincidences in the mass spectrum. For resonant freezeout, the resonance mass $m_a$ must be within a percent of twice the dark matter mass, a striking coincidence.  For cascade decays, the intermediate state(s) $\\phi$ must be within again about a percent of the dark matter mass in order to have a sharp enough spectral feature to fit the data well.  These near-degeneracies are suggestive of a dark sector with one scale $\\Lambda_m$ setting the overall mass scale, and another $\\Lambda_G \\ll \\Lambda_m$ determining splittings.  Composite dark sectors are an appealing avenue to flesh out the models we have considered. Near-degeneracies in the dark sector spectrum, as necessary for the cascade annihilation model, are readily accommodated in composite sectors. Resonances above threshold pose a somewhat more complicated picture. As we know from heavy quarkonia in the SM, a spectrum with resonances slightly above threshold could certainly be obtained; the complication is that such resonances would necessarily be accompanied by a resonance {\\em below} threshold, giving the thermal relic abundance a detailed dependence on the parameters of the model.  We leave detailed model building to future work.       \\subsection*"
    },
    "1208/1208.3499_arXiv.txt": {
        "abstract": "We confirm \\nconfirm\\ planets in \\nsystems\\ planetary systems by showing the existence of statistically significant anti-correlated transit timing variations (TTVs), which demonstrates that the planet candidates are in the same system, and long-term dynamical stability, which places limits on the masses of the candidates---showing that they are planetary.  % All of these newly confirmed planetary systems have orbital periods that place them near first-order mean motion resonances (MMRs), including 6 systems near the 2:1 MMR, 5 near 3:2, and one each near 4:3, 5:4, and 6:5.  In addition, several unconfirmed planet candidates exist in some systems (that cannot be confirmed with this method at this time).  A few of these candidates would also be near first order MMRs with either the confirmed planets or with other candidates.  One system of particular interest, \\kepotfo\\ (\\koiotfo ), is a pair of planets orbiting a 12th magnitude, giant star with radius over three times that of the Sun and effective temperature of 4900~K---among the largest stars known to host a transiting exoplanetary system. ",
        "introduction": "NASA's \\kepler\\ mission continues to identify many candidate transiting exoplanet systems, which now number nearly 2000 \\citep{Borucki:2011,Batalha:2012}.  The process of confirming or validating these planet candidates as real exoplanets often requires a significant amount of analysis and effort; and consequently the number of confirmed planets lags far behind the new candidate discoveries.  In an effort to ameliorate this situation we published a new approach to planet confirmation that requires a less detailed analysis \\citep{Fabrycky:2012a,Ford:2012a,Steffen:2012a} and can generally be accomplished in much less time.  This approach essentially relies on demonstrating that two transiting candidates are in the same system and that their masses are planetary.  Both aspects of this confirmation method rely on dynamical interactions.  First, planets are shown to be in the same system by looking for anticorrelated transit timing variations (TTVs) that arise from short-term changes in obital period due to planet-planet interactions within the system \\citep{Agol:2005,Holman:2005}.  Second, the candidates are shown to have planetary mass by requiring that the system be dynamically stable.  This method is a particularly useful means of identifying true planetary systems from among the false positive systems---allowing ground-based follow-up resources to be devoted to studying systems that are both real and dynamically interesting (as manifest by their significant TTV signal).  While detailed TTV analyses have been instrumental in confirming several multiplanet systems (e.g., Kepler 9 \\citep{Holman:2010}, Kepler 11 \\citep{Lissauer:2011a}, Kepler 18 \\citep{Cochran:2011}, and Kepler 36 \\citep{Carter:2012}), 21 planets in 10 multiplanet systems have been confirmed by these new methods that we apply again here (Kepler systems 23 through 32, \\citet{Fabrycky:2012a,Ford:2012a,Steffen:2012a}).  These 10 confirmed multiplanet systems include several that have particularly interesting properties and are likely to be the subject of future investigations.  For example: all of the systems are near mean-motion resonance (MMR), Kepler-25 (KOI-244, \\citet{Steffen:2012a}) is relatively bright at Kp $\\sim 10$, Kepler-30 (KOI-806, \\citet{Fabrycky:2012a}) has TTVs that deviate from a constant period by nearly a day over the course of the \\kepler\\ data, and many (Kepler 23, 24, 26, 31, and 32) have additional planet candidates that may be confirmed or validated.  As more data from \\kepler\\ are gathered, new systems that can be confirmed by the means outlined above can be found.  In this paper we apply the same methods used in \\citep{Fabrycky:2012a,Ford:2012a,Steffen:2012a} to 2 additional quarters of \\kepler\\ data (all data through Q8) and confirm \\nconfirm\\ new planets in \\nsystems\\ systems.  This paper is organized as follows.  In \\S \\ref{secProperties} we summarize the stellar properties for the host stars in these systems and the basic orbital and physical characteristics of the planets.  The dynamical confirmation of the relevant \\nconfirm\\ planets is presented in \\S \\ref{secConfirm}.  % We make concluding remarks in \\S \\ref{secConclusion}.  Data and results from some of the analyses performed in this paper are shown in the Appendix. ",
        "conclusions": "All of the systems that we confirm here and in similar previous studies are near first-order MMR.  In this case there are six pairs near the 2:1 MMR (\\kepofe , \\kepstz , \\kepssf , \\kepetn , \\kepotfo , and \\kepotsz ), five pairs near the 3:2 MMR (\\keptfe , \\kepees , \\kepnzf , \\kepotts , and \\kepoftn ), and one pair each in the 4:3 (\\keptzes\\ c and d ), 5:4 (\\keptzes\\ b and c), and 6:5 (\\keptst ).  Several systems show additional planet candidates that, due to their presence in known multi-transiting systems, are very likely to be true planets as well \\citep{Lissauer:2012}.  If we consider all of the KOIs to be planets, then a few systems will have multiple pairs of planets near first order MMRs with some of these pairs being adjacent links in a near resonant chain including: \\kepstz\\ near a 3:2 -- 2:1 chain (3:2:1), \\kepetn\\ near a 2:1 -- 2:1 chain (4:2:1), \\kepnzf\\ with one pair near the 2:1 and a separate pair near the 3:2, and \\keptzes\\ near a 5:4 -- 4:3 chain (20:15:12).  The balance of the planet pairs in these systems, if near a resonance at all, must be near a resonance of higher order.  These planetary systems, like the other systems that have been confirmed via TTV analyses, are of particular interest for long-term scrutiny with follow-up observations and dynamical studies.  The detailed dynamics of planetary systems yield meaningful information about the evolutionary histories of the orbital architectures of the system.  For example, the differences in the fraction of systems that show multiple candidates and detectabe TTVs among smaller planet candidates with few-day orbital periods compared with the larger, Jupiter-sized candidates in the same period range---these hot Jupiters have no compelling signs of current planet-planet dynamics---gives strong clues as to their likely distinct dynamical history compared to the bulk of the exoplanet population \\citep{Steffen:2012b}.  These newly confirmed planetary systems continue to show the value that transit timing variations have and will have in transiting planet endeavors.  As new planets are found, with increasingly smaller sizes and longer orbital periods, TTVs provide the only viable means of determining planetary masses.  TTVs, along with statistical validation via the techniques employed in \\citep{Torres:2011} and \\citep{Lissauer:2012}, will likely be the primary means to demonstrate the planetary nature of the planet candidates that the \\kepler\\ spacecraft identifies (in part because the stars in the \\kepler\\ field are often too dim to make good Radial Velocity targets).  Two of these systems show a modest gap in the orbital period ratios of adjacent pairs of KOIs.  If planetary systems tend to be dynamically packed (e.g. \\citet{Barnes:2004,Lissauer:1995,Laskar:2000}) then there may be additional undetected planets in these gaps.  In this set of systems, the most promising examples for future investigations include \\kepofe\\ which has a gap that is over a factor of 4 in period ratio and \\koitfe\\ which has a gap that is nearly a factor of three in orbital period.  A detailed TTV analysis of these two systems would provide useful constraints on or support for the hypothesis of dynamically packed systems.  The planetary system orbiting the giant star \\kepotfo\\ is another interesting individual system---being an early discovery of a transiting multiplanet system hosted by a giant.  Chromospheric emission from giant stars causes the star to be limb brightened in certain pass bands (e.g. the Calcium H and K lines).  Comparison of the transit duration in these bands with the duration from the broad band Kepler data may yield a direct measurement of the size of the chromosphere \\citep{Assef:2009}.  Continued monitoring by the \\kepler\\ spacecraft and future study will likely yield more precise mass measurements of the planets in these systems and consequently provide constraints on their bulk composition.  Such mass measurements will in turn improve our understanding of planet formation processes outside of our own solar system and are useful to place or solar system in the context of the general planet population."
    },
    "1208/1208.6275_arXiv.txt": {
        "abstract": "{ The Pierre Auger Observatory detects the highest energy cosmic rays. Calorimetric measurements of extensive air showers induced by cosmic rays are performed with a fluorescence detector. Thus, one of the main challenges is the atmospheric monitoring, especially for aerosols in suspension in the atmosphere. Several methods are described which have been developed to measure the aerosol optical depth profile and aerosol phase function, using lasers and other light sources as recorded by the fluorescence detector. The origin of atmospheric aerosols traveling through the Auger site is also presented, highlighting the effect of surrounding areas to atmospheric properties. In the aim to extend the Pierre Auger Observatory to an atmospheric research platform, a discussion about a collaborative project is presented. \\PACS{ {95.85.Ry}{cosmic rays}   \\and {95.55.Cs}{ground-based ultraviolet, optical and infrared telescopes} \\and {92.60.e}{atmospherics} \\and {92.20.Bk}{aerosols} \\and {92.60.-e}{properties and dynamics of the atmosphere} } % ",
        "introduction": "\\label{sec:intro} The Pierre Auger Observatory~\\cite{PAO_1,PAO_2} is the largest operating cosmic ray observatory ever built. It was designed to measure the flux, arrival directions and mass composition of cosmic rays from $10^{18}~$eV (electron-volt) to the very highest energies. When cosmic rays enter the atmosphere, they induce extensive air showers composed of secondary particles. Charged particles excite atmospheric nitrogen molecules, and these molecules then emit fluorescence light in the $300-400~$nm range~\\cite{AIRFLY,ArquerosFY}. At the Pierre Auger Observatory, the atmosphere is used as a giant calorimeter, representing a detector volume larger than $30\\,000~$km$^3$. To minimise as much as possible the systematic errors of the fluorescence measurements, atmosphere properties have to be continuously monitored~\\cite{AugerATMON}. During the development of an extensive air shower, the production rate of fluorescence photons depends on the temperature, pressure and humidity of the air~\\cite{BiancaFY,VazquezFY,DelphinePHIL}. Then, from their production point to the telescope, these photons can be scattered by molecules (by Rayleigh scattering) and/or atmospheric aerosols (by Mie scattering). Thus, to track atmospheric parameters, an extensive atmospheric monitoring system has been developed that covers the whole array.  \\begin{figure}[t]% \\centering \\resizebox{0.65\\textwidth}{!}{% \\includegraphics{fig1.pdf} } \\caption {{\\bf Atmospheric monitoring map of the Pierre Auger Observatory (from~\\cite{MyICRC}).} Gray dots show the positions of surface detector stations (SD). Black segments indicate the fields of view of the fluorescence detectors (FD) which are located in four sites, called Los Leones (LL), Los Morados (LM), Loma Amarilla (LA) and Coihueco (CO), on the perimeter of the surface array. Each FD site hosts several atmospheric monitoring facilities.} \\label{fig:MonitorArray} \\end{figure}  The different experimental facilities and their locations are shown in Fig.~\\ref{fig:MonitorArray}. Atmospheric properties at ground level are provided by a network of five weather stations located at each fluorescence detector (FD) site and at the Central Laser Facility (CLF). They furnish atmospheric state variable measurements every five minutes. Also, meteorological radio-sonde flights with balloons have been operated (for more details, see~\\cite{FocusPoint:BiancaMartin}). For the aerosol component, the central laser facility fires $50~$vertical shots every $15~$minutes during FD operations. Fluorescence telescopes, recording the Ultra-Violet (UV) laser tracks, are able to deduce the aerosol optical depth at different altitudes. In November 2008, a second laser facility called the eXtreme Laser Facility, or XLF, was deployed. These aerosol profile measurements are complemented by cloud measurements by four elastic backscattering lidars located at each eye (for more details, see~\\cite{FocusPoint:lidar}). A Raman lidar currently under test in Colorado (USA) is scheduled to be moved to the Auger Observatory for the Super-Test-Beam project~\\cite{WienckeICRC}. To improve our knowledge of photon scattering on aerosols, two Aerosol Phase Function monitors (APF) have been installed at the Coihueco and Los Morados FD sites. The APF instruments generate a collimated horizontal light beam produced by a Xenon flasher. The light passes in front of one FD site. The aerosol attenuation depends on the incident wavelength. This measurement is the main goal of two optical telescopes in Auger, the Horizontal Attenuation Monitor, or HAM~\\cite{HAM_proceeding}, and the (F/ph)otometric Robotic Telescope for Atmospheric Monitoring, or FRAM~\\cite{FRAM_proceeding}. % ",
        "conclusions": ""
    },
    "1208/1208.1380_arXiv.txt": {
        "abstract": "Galaxies become red and dead when the central supermassive black hole (SMBH) becomes massive enough to drive an outflow beyond the virial radius of the halo. We show that this final SMBH mass is larger than the final SMBH mass in the bulge of a spiral galaxy by up to an order of magnitude. The $M - \\sigma$ relations in the two galaxy types are almost parallel ($M \\propto \\sigma^{4+\\beta}$, with $\\beta < 1$) but offset in normalization, with the extra SMBH mass supplied by the major merger transforming the galaxy into an elliptical, or by mass gain in a galaxy cluster.  This agrees with recent findings that SMBH in two Brightest Cluster Galaxies are $\\sim 10\\times$ the expected $M-\\sigma$ mass. We show that these results do not strongly depend on the assumed profile of the dark matter halo, so analytic estimates found for an isothermal potential are approximately valid in all realistic cases.  Our results imply that there are in practice actually {\\it three} $M - \\sigma$ relations, corresponding to spiral galaxies with evolved bulges, field elliptical galaxies and cluster centre elliptical galaxies. A fourth relation, corresponding to cluster spiral galaxies, is also possible, but such galaxies are expected to be rare. All these relations have the form $M_{\\rm BH} = C_n\\sigma^4$, with only slight difference in slope between field and cluster galaxies, but with slightly different coefficients $C_n$.  Conflating data from galaxies of different types and fitting a single relation to them tends to produce a higher power of $\\sigma$. ",
        "introduction": "In the past decade, observations of large numbers of galaxies have revealed important correlations between the properties of galaxy spheroids and the supermassive black holes (SMBHs) they host. Two of these relations are particularly interesting: the black hole mass -- host spheroid velocity dispersion ($M - \\sigma$) relation \\citep{Ferrarese2006ApJ} and the black hole - bulge mass ($M - M_{\\rm b}$) relation \\citep{Haering2004ApJ}. Together, they strongly imply that SMBHs coevolve with their hosts, and probably affect each other's evolution through some form of feedback.   There have been numerous attempts to explain these relations, both analytically \\citep{Silk1998A&A, King2003ApJ, King2005ApJ, Murray2005ApJ, Power2011MNRAS} and numerically \\citep[e.g.][]{DiMatteo2005Natur, Booth2009MNRAS}. The AGN wind feedback model \\citep[and references therein]{King2003ApJ, King2005ApJ, King2010MNRASa, Zubovas2012ApJ} is particularly promising, as it explains not only the observed correlations, but also other observable properties of galactic winds \\citep[e.g.][]{Pounds2003MNRASb, Pounds2003MNRASa,Pounds2011MNRAS, Tombesi2010A&A, Tombesi2010ApJ} and outflows\\footnote{For clarity, we distinguish between the black hole {\\it wind}, the gas coming directly from the accretion disc around the SMBH, and the {\\it outflow}, the outward movement of swept--up interstellar gas from the host galaxy. These two components undergo respectively reverse and forward shocks on each side of the contact discontinuity separating them. See the Figure in \\citet{Zubovas2012ApJ} for more details} \\citep[e.g.][]{Rupke2011ApJ, Sturm2011ApJ}. In this picture, the wind and outflow shocks initially occur fairly close to the SMBH, and are efficiently Compton--cooled. The outflow retains only the ram pressure $\\simeq L_{\\rm Edd}/c$ of the black hole wind and has kinetic energy lower by a factor $\\sigma/c \\sim 10^{-3}$. Only when the SMBH mass approaches the critical mass $M_{\\sigma} \\simeq 3.7 \\times 10^8 \\; \\msun$ (see equation \\ref{eq:msigma} below) do the shocks move further away. The resulting geometrical dilution of the quasar radiation field now makes Compton cooling inefficient. The shocks are no longer cooled, and the outflows become energy--driven.  In two recent papers \\citep{King2011MNRAS, Zubovas2012ApJ} we investigated the properties of large (kiloparsec)--scale outflows in galaxies.  These must be energy--driven, i.e. the outflow kinetic energy rate $\\dot{E}_{\\rm out} = \\dot{M}_{\\rm out} v_{\\rm out}^2 / 2 \\simeq 0.05 L_{\\rm Edd}$, where $L_{\\rm Edd}$ is the Eddington luminosity of the driving quasar.  We showed that energy--driven outflows are able to clear galaxies of gas, turning them into red--and--dead spheroids.  However, in these papers, we only briefly touched upon the question of how much the SMBH mass grows as it expels the outflow from the galaxy, and this remains uncertain. There are two discrepant claims: \\citet{Power2011MNRAS} used an energy argument to conclude that the SMBH only needs to grow its mass by $\\sim40\\%$ to expel the surrounding gas past the virial radius, while in \\citet{King2011MNRAS} we found that an outburst duration of $\\sim 10^8$~yr is required, which would lead to the SMBH mass growing by about an order of magnitude.  We return to this problem in this paper. We also consider the difference between evolved spiral and elliptical galaxies in more general terms, as well as the dependence of the outflow properties on the galaxy environment. We find only a small uncertainty in our results stemming from the assumption of an idealized background potential (Section \\ref{sec:propagation}). By considering the distribution of energy in various components of the outflow we resolve the discrepancy between the two claims (Section \\ref{sec:outflows}). We conclude that SMBHs in elliptical galaxies should grow to a mass a few times higher than in spirals for a given value of $\\sigma$. We consider the effects of gas depletion and replenishment in a galaxy and find that galaxies in cluster environments should have slightly more massive SMBHs than their counterparts in the field (Section \\ref{sec:clusters}). We summarize and discuss our results in Section \\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth]{msig_pred_cor.eps} \\caption{The four $M-\\sigma$ relations (solid lines) and their combined effect on observational fits (dashed line). All solid lines have slopes $M \\propto \\sigma^4$ and the dashed line has $M \\propto \\sigma^6$. The grey area is the approximate locus of data points in Fig. 3 of \\citet{McConnell2011Natur}.} \\label{fig:msigdiag} \\end{figure}    We have shown that supermassive black holes residing in the bulges of spiral galaxies should have an $M-\\sigma$ relation that can be approximated by a power law slightly steeper than $M \\propto \\sigma^4$, if gas depletion due to star formation is considered. Typically, such galaxies would have $M < M_\\sigma$ as predicted in our earlier papers \\citep{King2003ApJ,King2010MNRASa,King2011MNRAS}. During an event triggering strong central accretion, such as a major merger, the SMBH is likely to grow by a further factor of $\\lesssim 7.5$ (i.e. for $\\lesssim 2$ Salpeter times) before it can drive the gas out of the galaxy halo. This factor is an upper limit, because we calculated it for an isothermal background potential, which is the most difficult to escape from.  One might think that since the SMBH is not immediately aware of the processes occurring around the virial radius of the halo, it could grow for longer than the calculated $\\sim 2$ Salpeter times, as long as the gas reservoir feeding it is not depleted. However, the total mass of gas that must be accreted in this process is of order $10^9 \\msun$, perhaps even more. This is too high to reside in a single accretion disc, and the SMBH is very probably fed in a series of stochastic feeding events \\citep{King2007MNRAS, Nayakshin2007arXiv}. Feedback therefore affects the SMBH feeding reservoir, and so it seems more likely that the reservoir is fairly efficiently depleted. While gas may fall back on to the SMBH if it has not yet escaped the halo and trigger subsequent bursts of activity, it cannot do so once the halo is depleted, so the AGN should switch off very soon after the outflow clears the galaxy.   Our results suggest that there should be a difference between the $M - \\sigma$ relations for spiral star-forming galaxies and for red and dead ellipticals. The latter should have SMBHs that are systematically overmassive when compared with the prediction of the $M - \\sigma$ relation for the spiral galaxies. This mass difference is not huge - less than an order of magnitude - but it may become apparent when better SMBH mass and spheroid velocity dispersion measurements become available.    There is another possible offset in the $M - \\sigma$ relation depending on the galaxy environment. Galaxies in clusters, especially those residing near cluster centres, may accrete gas from the hot cluster halo via cooling flows. This process may balance the depletion by star formation at a gas fraction $\\sim 0.2$, preserving the old $M - \\sigma$ relationship (eq. \\ref{eq:msigma}) in galaxies that had been spirals by the time they entered the cluster environment; even though they lose most of the gas from the outskirts, the central spheroidal component remains unperturbed and the analysis in Section \\ref{ref:spiralbulge} still applies. Cluster ellipticals would also be a factor $\\lesssim 7.5$ more massive. On average, cluster galaxies have black holes slightly more massive (by a factor $\\sim 2.5$) than their field counterparts. Our findings are consistent with the recent discovery of two SMBHs in the centres of giant elliptical galaxies that have masses in excess of $10^{10} \\; \\msun$. They are both Brightest Cluster Galaxies and our model predicts that their masses should be approximately an order of magnitude above the relation derived in earlier works.   In summary, we have found that there may be three (or possibly even four) $M-\\sigma$ relations in total, depending on galaxy morphology and environment:  \\begin{itemize} \\item spiral galaxies with evolved bulges have $M_{\\rm BH} \\sim 1.3 \\times 10^8 \\sigma_{200}^{4+\\beta} \\msun$ with $\\beta \\lesssim 1$ depending on star formation efficiency and the relation between the bulge size and $\\sigma$; \\item spiral galaxies with evolved bulges residing in gas-rich cluster environments have $M_{\\rm BH} \\sim 3.7 \\times 10^8 \\sigma_{200}^4 \\msun$ (the original $M-\\sigma$ relation); such galaxies, however, may be extremely rare due to high merger probability in clusters; \\item field elliptical galaxies have $M_{\\rm BH} \\sim 9.8 \\times 10^8 \\sigma_{200}^{4+\\zeta} \\msun$; \\item elliptical galaxies close to cluster centres have $M_{\\rm BH} \\sim 28 \\times 10^8 \\sigma_{200}^4 \\msun$. \\end{itemize}  There is some uncertainty in the slope, especially for field galaxies. It is clear that, as bigger bulges have larger velocity dispersions, the slope should be steeper than $4$, but its precise value depends strongly on the assumed relation between $R_{\\rm b}$ and $\\sigma$ (both its slope and intercept).  The differences among the four relations are small, at most factors of a few in the normalization. Intrinsic scatter in each population may blur the distinctions further. In addition, since elliptical galaxies tend to have higher velocity dispersions than spirals, and cluster galaxies than field galaxies, the overall effect of combining the relations may be to increase the observed slope of the $M-\\sigma$ relation. We illustrate this in Figure \\ref{fig:msigdiag}. We plot the four $M-\\sigma$ relations (assuming $\\beta = 0$) over a polygon representing the approximate locus of data points taken from Figure 3 of \\citet{McConnell2011Natur}. All four lines intersect the locus. The large fraction of the locus being below the four lines is reasonable since the predicted relations are only upper limits for different environments. Furthermore, there is a hint that higher values of $\\sigma$ correspond to higher $M-\\sigma$ relations. We also plot a line with $M \\propto \\sigma^6$, which roughly bisects the observed locus and so might emerge from a fit to the data unsorted by galaxy type, especially in the $\\sigma > 100$~km/s part, where we do not expect contamination from nuclear stellar clusters \\citep{McLaughlin2006ApJ, Ferrarese2006ApJ}.   There have been recent tentative claims that the $M-\\sigma$ relation may be slightly different for different galaxy types, and that the general relation may show an upturn at higher values of $\\sigma$ \\citep{Marconi2003ApJ,McConnell2011Natur}. Our considerations here suggest that both features may be natural consequences of wind feedback."
    },
    "1208/1208.3666_arXiv.txt": {
        "abstract": "We demonstrate that a quasi-uniform cosmological seed field is a much less suitable seed for a galactic dynamo than has often been believed. The age of the Universe is insufficient for a conventional galactic dynamo to generate a contemporary galactic magnetic field starting from such a seed, accepting conventional estimates for physical quantities. We discuss modifications to the scenario for the evolution of galactic magnetic fields implied by this result. We also consider briefly the implications of a dynamo number that is significantly larger than that given by conventional estimates. ",
        "introduction": "We start from the concept that galactic magnetic fields are  the result of dynamo action driven by galactic differential rotation and mirror asymmetric interstellar turbulence. This immediately raises the question of the origin and nature of the seed field from which the mature field arises. Two options are apparent.  A very weak magnetic field can be created in a protogalaxy by a battery effect, and then excited by turbulent motions (the small scale dynamo), to produce a small-scale field of strength up to approximate equipartition with the turbulent motions. This then serves as a seed for a mean-field galactic dynamo, e.g. \\cite{b96,r06}. This scenario appears uncontroversial given our present understanding of dynamo theory and the physics of galaxies.   Another possibility is that the seed field for the galactic dynamo has survived from a very early time in the Universe, see for example the review by \\cite{ss05}. This picture also looks attractive, especially given recent claims that a significant magnetic field exists in the intergalactic medium (\\cite{ns09}); such a field has been discussed earlier in various other contexts, see e.g. \\cite{kz08}. This scenario raises the question of how to form such a magnetic field of suitable scale, and then to maintain it against Ohmic losses between the epochs of recombination and galaxy formation. The most straightforward possibility perhaps is to consider a strictly homogeneous magnetic field which exists from the epoch of the Big Bang, see \\cite{z65}. It has been widely accepted that a galactic mean-field dynamo can, without any obvious problems, use this field (provided it is sufficiently strong) as a seed from which to produce contemporary galactic magnetic fields. This idea has been accepted as obvious in many reviews, e.g. \\cite{b96}.  Observational tests to distinguish between these possibilities are difficult if not impossible currently. However it may be feasible to formulate observable predictions from each of them, for confrontation  with future observations with forthcoming telescopes such as LOFAR and SKA. Some predictions for tests with  LOFAR were already suggested in \\cite{a09}, \\cite{metal2011}.  Our original intention when starting this work was to suggest some observational tests for the second scenario. In the spirit of this intention we performed numerical simulations of a galactic dynamo model with a homogeneous seed. Contrary to the usual expectations, we arrived at the conclusion that a homogeneous seed may not be as suitable as has been widely thought for generation of a contemporary galactic magnetic field. The aim of this paper is to present and develop this finding.  A preliminary description of this work appears in Sokoloff \\& Moss (2011); here we elaborate and expand our investigation.  \\begin{figure*} \\begin{center} \\begin{tabular}{ll} \\label{t_evol} (a)\\includegraphics[width=0.44\\textwidth]{ts4+4x.eps} & (b)\\includegraphics[width=0.44\\textwidth]{tsr2.eps} \\\\ \\end{tabular} \\end{center} \\caption{The evolution of the magnetic field for $R_\\alpha=0.8, R_\\omega=8$: (a) uniform seed field with $B_x(t=0)=1$, alpha-quenched shown as solid, $B_x(t=0)=10^{-6}$, linear by the broken curve; (b; adapted from \\cite{sm11}) random seed field, $B{\\rm r.m.s}(t=0)=1$, alpha-quenched calculation } \\end{figure*} ",
        "conclusions": "We see that we have come to a rather unexpected conclusion: a cosmological magnetic field that is uniform on the scale of the young galaxy is a \"bad\" seed for mean-field galactic dynamos.  Of course, this conclusion holds for the conventional estimate for the intensity of dynamo action, as measured by the dynamo number $D$, and estimated using knowledge of contemporary galaxies. It is possible that this estimate should be reconsidered for very young galaxies -- see e.g. the example illustrated in Fig.~3. In principle, this possibility might be pursued to attempt to produce bisymmetric magnetic configurations in galaxies. However, the situation is unclear, and galactic dynamos can anyway work without any cosmological seed field.  One possibility is that a relatively strong initial field (whose origin we are not considering) undergoes dynamical instabilities, which generate small-scale magnetic fields which contain nonvanishing $m=0$ components, and these then grow in a similar manner to that shown in Fig.~1b and the right hand column of Fig.~2.  Related in some ways to this is the possibility that a uniform field, being affected by, say, small-scale turbulence, acquires all modes including $m=0$. In either of these cases a mean-field dynamo can take as a seed the $m=0$ component so generated, and amplify it up to the currently observed field strengths and configurations. Thus our result does not exclude in principle the role of cosmological seed  fields. The point however is that the growth of the initial field component $m=0$ is then determined by a particular distortion that creates it, rather than being present in the cosmological field itself. Numerous such distortions will  occur during early galactic evolution, and so the real difference between two scenarios seems to be much less important than is usually assumed.  Our result is valid for both strong and weak initial fields, and the inclusion of a simple saturation mechanism makes essentially no difference; with a strong seed  after a short initial period the evolution is essentially linear.  Our interpretation of the results obtained needs however some comments. The no-$z$ model for a strictly homogeneous seed (i.e. uniform, in any direction) describes the magnetic field component perpendicular to the rotation axis while the component parallel to the rotation axis remains unaffected by the rotation.  If there is a large-scale homogeneous seed of cosmological (or very early pre-galaxy formation) origin then the component parallel to $\\Omega$ is of aligned dipolar parity, and the perpendicular component has the parity of a perpendicular dipole (i.e. $m=1$). The no-$z$ model can handle the latter case (but not the former, being explicitly restricted to fields of even parity with respect to the galaxy mid-plane). It shows that it is a \"bad\" seed field in sense that corresponding dynamo generated magnetic field develops too slowly to generate the present-day galactic magnetic field. The  code can also handle a axisymmetric quadrupole-like seed, which grows normally (seed is, e.g., purely radial in the no-$z$ formulation). Classical axisymmetric disc dynamo models show that axisymmetric  fields of quadrupolar parity are preferred over those of dipolar parity, e.g. \\cite{rss88} and much subsequent work. Thus we deduce that large scale homogeneous seeds are inefficient, whatever their orientation -- large-scale quasi-homogeneous seed fields can hardly have quadrupolar parity.    In this context, another limitation of the no-$z$ model to be considered is that it cannot explicitly include the effects of galactic winds blowing out of the disc, or of turbulent diamagnetism (in some ways these have opposing effects). Both of these mechanisms were investigated by Brandenburg {\\it et al.} (1993), who found that solutions with quadrupolar parity remained preferred -- consistent with a basic assumption of the no-$z$ model. \\cite{metal10} similarly demonstrated that the preference for quadrupolar fields persisted under a wide range of conditions; globally mixed parity solutions are also possible, although these are usually dominantly quadrupolar in the disc region.  We recognize that widely discussed mechanisms for magnetic field generation in the early universe give rise to small-scale fields. In such cases, the subsequent evolution of galactic fields can not expect to be distinguished from that following small-scale dynamo action in the proto-galaxy. In this paper we have mainly considered the consequences of an alternative scenario (which we do not particularly advocate), that a large-scale quasi-homogeneous cosmological field is present at the time of galaxy formation.  Note also that contemporary dynamo theory considers dynamo models that are much more detailed than is used by our simple mean-field dynamo model, e.g. \\cite{hetal09}, \\cite{kd11}, \\cite{getal08}. Models based on direct numerical simulations of the induction equation rather than on mean-field dynamo equations, include however both the galactic dynamo action and distortion of the uniform field by small-scale motions simultaneously, and demonstrate the joint action of both effects. In contrast,  in this paper we have separated them, in order to elucidate the basic principles.  We are grateful to our colleagues in research on magnetic field evolution in very young galaxies, T. Arshakian, R. Beck, M. Krause, R. Stepanov, for fruitful discussions which stimulated our writing this the paper. DS acknowledges financial support from NORDITA for his participation in the meeting."
    },
    "1208/1208.2276_arXiv.txt": {
        "abstract": "Planets and other low-mass binary companions to stars face a variety of potential fates as their host stars move off the main sequence and grow to subgiants and giants.  Stellar mass loss tends to make orbits expand, and tidal torques tend to make orbits shrink, sometimes to the point that a companion is directly engulfed by its primary. Furthermore, once engulfed, the ensuing common envelope (CE) phase can result in the companion becoming fully incorporated in the primary's envelope; or, if the companion is massive enough, it can transfer enough energy to eject the envelope and remain parked in a tight orbit around the white dwarf core.  Therefore, ordinary binary evolution ought to lead to two predominant populations of planets around white dwarfs: those that have been through a CE phase and are in short-period orbits, and those that have entirely avoided the CE and are in long-period orbits. ",
        "introduction": "\\label{sec:intro}  Many intermediate-mass stars have either stellar companions or brown dwarf or planetary companions \\citep{duquennoy+mayor1991, raghavan_et_al2010, schneider_et_al2011, wright_et_al2011}.  These systems can sometimes remain dynamically quiet for hundreds of millions or billions of years, and then undergo relatively rapid changes when the primary star evolves off the main sequence.  Various recent works have investigated the post-main-sequence fates of two-body systems \\citep{villaver+livio2007, carlberg_et_al2009, villaver+livio2009, nordhaus_et_al2010} and of many-body systems \\citep{veras_et_al2011, veras_et_al2012, kratter+perets2012}.  The present work is generally concerned with the former type of system --- i.e., the evolution of a single star and a single low-mass companion.  The status of planetary systems around post-main-sequence stars is starting to come into focus.  \\citet{zuckerman_et_al2010} find evidence of heavy-element atmospheric pollution in $\\sim$1/3 of DB white dwarfs, which they interpret as due to accretion of tidally shredded asteroids that were presumably scattered onto high-eccentricity orbits by distant planets.  \\citet{maxted_et_al2006} found evidence of perhaps the best-characterized substellar companion to a white dwarf, via high-resolution spectroscopy that revealed radial velocity variations with period 2~hours with semiamplitudes of $\\sim$28~km~s$^{-1}$ in the primary and $\\sim$188~km~s~$^{-1}$ in the secondary.  They infer a secondary mass of $\\sim$50~$M_J$ (where $M_J$ is the mass of Jupiter) --- clearly above the canonical ``planet'' mass but less massive than the hydrogen-burning mass limit \\citep{burrows_et_al2001}.  Provocative evidence of very close low-mass companions to an evolved star was reported by \\citet{charpinet_et_al2011}, who claim to find two sub-Earth-radius planets around the subdwarf star KIC 05807616 that are both orbiting inside 1.7~$R_\\sun$ (where $R_\\sun$ is the radius of the Sun).  If these planet candidates are actually planets, this discovery shows that exotic and theoretically unanticipated post-main-sequence planetary systems are possible.  \\citet{hogan_et_al2009} report a nondetection of warm companions around 23 nearby white dwarfs and suggest that $\\lesssim$5\\% of white dwarfs have companions with effective temperatures greater than 500~K between 60~AU and 200~AU in projected separation.  Several years ago, the pulsating white dwarf GD-66 was found to exhibit timing variations in pulsations that seemed to be consistent with light-travel-time delays caused by orbital motion around the common center of mass with a companion on a $\\sim$4.5-year orbit \\citep{mullally_et_al2007, mullally_et_al2008, mullally_et_al2009}.  However, recent high-precision astrometric observations of this system, an analysis of {\\it Spitzer} observations of the system, and follow-up observations by the team of original discovery, have complicated the planetary hypothesis for this system (\\citealt{farihi_et_al2012}; J. Hermes 2012, private communication). Finally, \\citet{johnson_et_al2011} has found a number of giant planets around slightly evolved, subgiant stars (see also \\citealt{sato_et_al2008, bowler_et_al2010}), some of which might be near the edge of engulfment or repulsion.  A number of lines of evidence, then, point to the existence of planets around evolved stars; even if some of the candidate systems eventually turn out not actually to be planets, it still seems clear that several tens of percent of white dwarfs have planets around them \\citep{zuckerman_et_al2010}.  It is worthwhile to recognize that the existence of planets orbiting post-main-sequence stars should be no surprise, given that that the very first planets discovered beyond our solar system were found around the millisecond pulsar PSR1257+12 \\citep{wolszczan+frail1992}. Whether planets around stellar remnants formed after the death of the primary, or were present during the main-sequence phase and survived stellar evolution, remains an open question.  It seems likely that both processes occur \\citep{hansen_et_al2009, tutukov+fedorova2012}; although there might be more reason to believe that planets could form around a pulsar than around a white dwarf, there is no consensus as yet.  Many jovian companions to main-sequence stars will become highly irradiated as their primaries evolve off the main sequence, thereby turning them into ``hot Jupiters,'' of sorts --- or red-giant hot Jupiters, as described by \\citet{spiegel+madhusudhan2012}.  The atmospheres of these companions might become transiently polluted by accretion both of the evolved star's wind and of dust and planetesimals \\citep{spiegel+madhusudhan2012, dong_et_al2010}.  Some of these planetary companions, or somewhat higher mass companions, will eventually be swallowed by their stars, which might contribute to the formation and morphology of planetary nebulae \\citep{nordhaus+blackman2006, nordhaus_et_al2007} and to the formation of highly magnetic white dwarfs \\citep{tout_et_al2008, nordhaus_et_al2011}, although there are other formation models in the literature as well \\citep{garcia-berro_et_al2012}.  The remainder of this document is structured as follows: In \\S\\ref{sec:tides}, I show that the increased moment of inertia of an evolved star can cause a companion that had been slowly moving outward due to tidal torques to instead rapidly plunge into the primary.  In \\S\\ref{sec:2pops}, I argue that simple binary evolution leads to two populations of companions to white dwarfs, with a large gap in period (or orbital radius) between them.  In \\S\\ref{sec:sdB}, I comment briefly on the potential sub-Earth-sized planets around KOI 55 found by \\citet{charpinet_et_al2011}.  Finally, in \\S\\ref{sec:conc}, I summarize and conclude. ",
        "conclusions": "\\label{sec:conc}  Post-main-sequence stellar evolution results in dramatic changes in the stellar radius and, therefore, in the orbits of companions to the stars.  Companions that are too close can either be directly swallowed by the expanding star or tidally dragged into merging with the star. The ensuing common envelope poses severe risks to the survival of the companion; massive enough companions can eventually transfer enough orbital energy to unbind the star, resulting in a tight post-CE orbit, while less massive companions continue spiraling inwards until they are tidally shredded and merge with their host stars.  Companions that are on distant enough orbits to avoid ever merging with their host stars move outward due to mass loss from the primary.  Nearly all objects that are massive enough to survive a CE have masses in excess of the deuterium-burning limit that is sometimes used to dilineate between planets and brown dwarfs \\citep{spiegel_et_al2011a}. For an object that is less than $\\sim$6 times the mass of Jupiter to end up inside 1~AU from a WD requires something more complex than simple binary/CE evolution.  In particular, forming the potentially habitable circum-white-dwarf terrsstrial objects considered by \\citet{agol2011} and \\citet{fossati_et_al2012} might require exotic circumstances involving multiple bodies, and the evolutionary paths to produce such worlds might be problematic for their subsequent habitability, as pointed out by \\citet{nordhaus+spiegel2013}.   \\vspace{1in}  {\\sc Acknowledgments}  This work was in part inspired by the ``Planets Around Stellar Remnants'' conference in Arecibo, Puerto Rico, January 23-27, 2012.  I thank Jason Nordhaus, John Johnson, Ruobing Dong, Scott Gaudi, Jeremy Goodman, and Piet Hut for illuminating conversations, and Alex Wolszczan for organizing the conference.  I gratefully acknowledge support from NSF grant AST-0807444 and the Keck Fellowship, and from the Friends of the Institute."
    },
    "1208/1208.2795_arXiv.txt": {
        "abstract": " ",
        "introduction": "Planetary atmospheres can be studied remotely using spectroscopy of the light that is reflected or the radiation that is emitted by the planet. Another method for studying planetary atmospheres is to measure the light of a bright source, such as a star or the Sun, as it is attenuated by (part of) the atmosphere. For Solar System objects, such transmission measurements are performed during stellar or solar occultations, when the planetary limb is in between the light source and the spacecraft or telescope. One advantage of transmission measurements is that the light source is often very bright, enabling high signal-to-noise measurements. Another advantage is that transmission measurements allow sampling of long path lengths through the atmosphere, since the light travels through the curved limb of the planet. Long path lengths increase the sensitivity to e.g.~trace gases with very small concentrations. Furthermore, altitudes in an atmosphere can be probed with a high vertical resolution, as only a very limited altitude range is probed by a single observation compared to e.g.~on-disc observations, and many altitudes can be sampled due to the high signal-to-noise. Also, occultation measurements rely on relative changes in the measured signal, and not on absolute signal levels, and are hence `self-calibrating' and less dependent on instrument drifts. As a result of these advantages, transmission measurements are very suitable for retrieving vertical profiles of trace gases. Disadvantages of solar and/or stellar occultations are that they require very special observing geometries, making it not trivial to target specific places on the planet. Furthermore, only a single location on the planet can be studied simultaneously, and the long path lengths prevent probing the lower, thicker layers of a planetary atmosphere. For solar occultations, the atmosphere can naturally only be studied at local twilight conditions. In recent years there have been many occultation measurements of Solar System planets. Besides the Earth, occultation transmission measurements have been performed of Venus \\citep[e.g.][]{van08,fed08,bel12}, Mars \\citep[e.g.][]{bla89,kra89,for09}, Jupiter \\citep{for03}, Saturn \\citep{kim12}, Titan \\citep{bel09,kim11}, and Pluto \\citep{hub88,ell03}, for which more are planned \\citep{ste08}. Besides measuring the atmospheric transmission, occultations can also be used to measure the atmospheric refraction, which we will not consider here.  Transmission measurements are one of the most valuable probes of the atmospheres of extrasolar planets (exoplanets). The transmission of an exoplanetary atmosphere can be measured when the orbit of the exoplanet with respect to the observer is such that the planet crosses the disc of its host star. By accurately measuring the wavelength-dependence of the decrease of star light during such a transit, a transmission spectrum of the planetary atmosphere can be derived \\citep{sea00,bro01,hub01}, although only the day-night boundary can be probed. These transit measurements have revolutionised our knowledge of exoplanet atmospheres and are now widely used to derive atmospheric properties of a number of transiting exoplanets. There are now several planets for which molecular absorption features have been identified through measuring their transits \\citep[e.g.][]{cha02,tin07,sne10,sea10}, as well as molecular and haze Rayleigh scattering \\citep[e.g.][]{lec08,pon08,sin11}. Also, there is an indication of high wind speeds on the planet HD 209458b \\citep{sne10}. The smallest planet for which there have been significant transit measurements is the `super-Earth' GJ 1214b \\citep{cha09}. Transit measurements allow the determination of the bulk composition of this planet \\citep{mil10}, but at present it is not clear whether it is has a water vapour atmosphere, or a hydrogen atmosphere with clouds \\citep{bea10,bea11,cro11,dem12}. Unfortunately, atmospheres with high molecular weight, like Earth, are hard to characterise in this way, as the spectral features are much reduced compared to hydrogen atmospheres. This is because the scale height of the atmosphere scales with the molecular weight.  The transmission of a planetary atmosphere is mainly determined by the extinction or total optical thickness of gases and particles, like haze and cloud particles, measured along the line of sight. The total optical thickness $\\tau_{\\rm t}$ can be split into an absorbing part, $\\tau_{\\rm a}$, and a scattering part, $\\tau_{\\rm s}$, using the single-scattering albedo $\\tilde{\\omega}$ of the mixture of gases and particles in the layer: \\begin{equation} \\tau_{\\rm t}= \\tilde{\\omega}\\tau_{\\rm t} + (1-\\tilde{\\omega}) \\tau_{\\rm t}= \\tau_{\\rm s} + \\tau_{\\rm a}. \\end{equation} The optical thicknesses $\\tau_{\\rm t}$, $\\tau_{\\rm a}$, and $\\tau_{\\rm s}$ and the single scattering albedo $\\tilde{\\omega}$ are usually wavelength dependent.  If the transmission $T$, i.e. the ratio of transmitted to incident flux, is purely determined by the extinction along the line of sight, it is simply given by the Beer-Lambert law: \\begin{equation} T= e^{-\\tau_{\\rm t}} \\label{eq.transm} \\end{equation} If the scattering optical thickness $\\tau_{\\rm s}$ along the line of sight is larger than zero, however, a fraction of the light that has been scattered out of the incident beam of light, will be scattered in exactly the forward direction after a single scattering or after multiple scattering events. This means that more light reaches the observer than expected from the extinction optical thickness of the atmosphere alone. In fact, large particles scatter light predominantly in the forward direction because of diffraction \\citep[see e.g.][]{han74}, possibly making the forward-scattered light contribute significantly to the total transmitted flux. Indeed, non-absorbing, perfectly forward-scattering particles would appear completely transparent in transmission, despite a non-zero optical thickness.  How a particle scatters the incident flux as a function of the scattering angle $\\Theta$ is described by the phase function $P(\\Theta)$. Figure~\\ref{fig.ssa+p0} shows the single-scattering albedo $\\tilde{\\omega}$ and the phase function in the forward direction, $P(0)$, of some particles found in atmospheres of Solar System planets. The phase functions have been normalised such that their integral over 4$\\pi$ steradians is unity. We show measured as well as calculated values of $P(0)$. The properties of Titan aerosol and martian dust particles are derived from measurements by \\citet{tom08} and \\citet{tom99}, respectively. The lines in Fig.~\\ref{fig.ssa+p0} are computed using Mie-theory \\citep[][]{der84}. Note that the strength of the forward-scattering part of the phase function is due to diffraction and as such depends mostly on the size of the particles, and less on their shape \\citep[e.g.][]{mis96}. Hence, Mie scattering should be accurate for the purpose of estimating the forward-scattering peak. For the Venus cloud particles, we used the optical constants of \\citet{pal75} and the size distributions for the two most dominant size modes from \\citet{gri93}. For the calculations for martian dust and ice cloud particles we used the optical parameters and size distributions as described by \\citet{kle09}. It can be seen that especially at short wavelengths $P(0)$, and hence the forward-scattering contribution to the transmission, can potentially be large, even though all these different particles are relatively small (with effective radii smaller than 2 $\\mu$m). The forward-scattering part of the phase function of larger particles, such as liquid water cloud particles found on Earth, can be much larger \\citep[see e.g. Fig.~11.13.6 in][]{mis06}.  \\begin{figure}[htp] \\centering \\includegraphics[width=0.9\\textwidth]{ssa+p0}  \\caption{Single-scattering albedos $\\tilde{\\omega}$ and phase functions in the forward direction, $P(0)$, of particles present in the atmospheres of Venus, Mars and Titan (see the text for details).} \\label{fig.ssa+p0} \\end{figure}   Analysing transmission measurements using Eq.~\\ref{eq.transm} and thus ignoring the (potential) contribution of forward-scattered flux to the transmission, which is how these measurements are usually analysed, will result in an error in the retrieved optical thickness of the planetary atmosphere. In this paper, we explore the contribution of the forward-scattered flux to the transmission signal using a three-dimensional Monte Carlo model, which simulates scattered light in a spherical, stratified atmosphere. Besides providing general results, we will also present more detailed results for Titan's atmosphere, to serve as a practical example. Finally, we draw conclusions regarding possible errors in retrievals from transmission measurements when scattering particles are present. ",
        "conclusions": "We have developed a three-dimensional, spherical Monte Carlo code to simulate the added contribution of scattered light to the flux that is transmitted through a planetary atmosphere during stellar or solar occultations and exoplanet transits. Our calculations with parameters of the exoplanet HD 189733b show that Rayleigh scattering does not add significantly to the transmitted flux that one would expect when the transmission depended solely on the extinction (absorption plus scattering) optical thickness of the atmosphere. However, forward-scattering particles can contribute significantly to the transmitted flux. For (slant) optical thicknesses larger than unity, the added contribution of forward-scattered flux can be as large or larger than the flux expected from the extinction alone. For exoplanet transits, this means that the magnitude of a transit of a planet with forward-scattering particles in its atmosphere can be several percent smaller than predicted when these particles are assumed to be fully absorbing. Such a difference in transit depth is easily comparable to the magnitude of measured and predicted spectral features in a transit spectrum.  With our Monte Carlo code, we performed detailed calculations of occultation spectra of Titan's atmosphere between 2.0 and 2.8 $\\mu$m and show that the haze abundance and trace gas mixing ratios are underestimated by $\\sim 8$\\% if the forward-scattering contribution to the transmitted flux is ignored in the retrieval algorithm. Also in such a case, the derived spectral slope of the haze particles is underestimated by a few percent. Figure~\\ref{fig.ssa+p0} shows that Titan's haze particles are fairly forward-scattering between 2~and 3$\\mu$m, but that cloud and haze particles can be at least an order of magnitude more forward-scattering at shorter wavelengths. This means that at shorter wavelengths, cloud, haze or gas abundances could be underestimated by several tens of percent if forward-scattering is ignored. Figure~\\ref{fig.ssa+p0} also shows that such strong forward-scattering is not uncommon in other atmospheres in the Solar System. Also the change of forward-scattering with wavelength can be much stronger at shorter wavelength, resulting in much larger errors in the derived spectral slope of forward-scattering particles for absorption-only models. Similarly, fast changes of the single-scattering albedo with wavelength can also cause significant errors in the derived spectral slope of cloud and haze particles.  We also compared our Monte Carlo calculations for three-dimensional, spherical planetary atmospheres to calculations for atmospheres composed of plane-parallel slabs, with the same optical thickness as the slant light paths through the spherical atmosphere, and with scattering properties weighted by the slant optical thickness of each atmospheric layer (see e.g.~\\citet{hub01,kim11} for examples).  For the cases we considered, we found these two calculations to match within 10\\% for optical thicknesses less than 10. This has the implication that the contribution of forward-scattered light to transmitted fluxes can be fairly well modelled using a plane-parallel assumption, for which radiative transfer routines are available that are much faster than a three-dimensional Monte Carlo code.   \\textbf{Acknowledgements}  We acknowledge financial support by the Netherlands Organisation for Scientific Research (NWO). We thank Andreas Macke for use of part of his code. We thank the two anonymous reviewers for their careful reading and useful suggestions.  \\clearpage"
    },
    "1208/1208.0335_arXiv.txt": {
        "abstract": "We determine the distribution of stellar surface densities, $\\Sigma$, from models of static and dynamically evolving star clusters with different morphologies, including both radially smooth and substructured clusters. We find that the $\\Sigma$ distribution is degenerate, in the sense that many different cluster morphologies (smooth or substructured) produce similar cumulative distributions. However, when used in tandem with a measure of structure, such as the $\\mathcal{Q}$-parameter, the current spatial and dynamical state of a star cluster can be inferred. The effect of cluster dynamics on the $\\Sigma$ distribution and the $\\mathcal{Q}$-parameter is investigated using $N$-body simulations and we find that, depending on the assumed initial conditions, the $\\Sigma$ distribution can rapidly evolve from high to low densities in less than 5\\,Myr. This suggests that the $\\Sigma$ distribution can only be used to assess the current density of a star forming region, and provides little information on its initial density. However, if the $\\Sigma$ distribution is used together with the $\\mathcal{Q}$--parameter, then information on the amount of substructure can be used as a proxy to infer the amount of dynamical evolution that has taken place. Substructure is erased quickly through dynamics, which can disrupt binary star systems and planets, as well as facilitate dynamical mass segregation. Therefore, dynamical processing in young star forming regions could still be significant, even without currently observed high densities. ",
        "introduction": "Stars form in clusters and associations  \\citep[e.g.][]{Carpenter00b,Lada03,Lada10}, and their spatial distribution appears to be smooth and continuous, i.e.\\,\\,there is no evidence of bi--modal star formation \\citep{Bressert10}. However, the morphologies of individual clusters appear to vary significantly \\citep{Cartwright04}, with some regions displaying a high amount of substructure (e.g.~Taurus and Chamaeleon) and others displaying a smooth, centrally concentrated morphology (e.g.~$\\rho$~Ophiuchus and IC\\,348). This suggests that either star formation in clusters is not universal and may depend on the local environment, or that all clusters form with the same morphology which is later altered by dynamical interactions.  There are several methods to look for structure in clusters, and even clusters themselves. The study by \\citet{Cartwright04} used a minimum spanning tree analysis, which has also been used to quantify the amount of mass segregation in clusters \\citep{Allison09a} and can be used to search for clusters, or ``clustering'' in crowded fields \\citep{Gutermuth09,Schmeja11}, in tandem with the mean separation between stars, to define the $\\mathcal{Q}$--parameter. The $\\mathcal{Q}$--parameter quantifies whether the cluster is substructured or radially concentrated, and to what extent, for a given morphology.  The stellar surface density distribution has also been used to study the structure of star forming regions. \\citet{Larson95} showed that the structure in Taurus was hierarchical, but broke down at the binary regime (corresponding to the local Jeans length), and \\citet{Simon97} found the same result for Ophiuchus and the Trapezium cluster. However, \\citet*{Bate98b} repeated this analysis for the Trapezium cluster and found that the break between random stars and binaries does not correspond to the Jeans length in this cluster.  Furthermore, \\citet{Kraus08} found that the break between random stars and binaries was 11\\,000\\,au in Upper Sco, compared to 17\\,000\\,au in Taurus, and suggested the break is a function of the age (and initial density) of the cluster, and is caused by dynamical evolution.  Recently, the distribution of stellar surface densities  ($\\Sigma$) was also used to evaluate the definition of what constitutes a star cluster. \\citet{Bressert10} examined a volume-limited sample of nearby star-forming regions and calculated the distribution of $\\Sigma$ in each region. \\citet{Bressert10} then compared the cumulative distribution of $\\Sigma$ for all the regions (apart from the central region of the Orion Nebular Cluster (ONC), where the determination of $\\Sigma$ was not possible for every star due to crowding) to various definitions of clustering, based on stellar surface density thresholds. Due to a lack of obvious bi-modality in the distribution, they concluded that stars in the local solar neighbourhood form in a continuous spatial distribution.  In this paper, we determine the effectiveness of the surface density distribution, $\\Sigma$, in tandem with the $\\mathcal{Q}$-parameter, for providing information on the dynamical state of star clusters. We begin by analysing static star clusters with various morphologies to look for variations in the shape of the $\\Sigma$ distribution. We then dynamically evolve radially smooth Plummer sphere clusters, and substructured fractal clusters, and determine the $\\Sigma$ distribution and the $\\mathcal{Q}$-parameter as a function of time. We describe the set-up of the clusters in Section~\\ref{method}; we present our results in Section~\\ref{results}; we provide a discussion in Section~\\ref{discuss} and we summarize and conclude in Section~\\ref{conclude}. ",
        "conclusions": "\\label{conclude}  In this paper, we have modelled both static and evolving star clusters with different morphologies, and determined the distribution of stellar surface densities, $\\Sigma$, and the $\\mathcal{Q}$-parameter. In our $N$-body simulations we determine the $\\Sigma$ distribution at 0, 1, 5 and 10\\,Myr. Our conclusions are the following:  (i) The distribution of $\\Sigma$ is degenerate. Many different morphologies reproduce a smooth, continuous distribution of stellar densities. We find that a substructured star forming region (either a fractal, or an association) has a narrower cumulative distribution than a centrally concentrated spherical cluster, such as a Plummer sphere or a King profile. However, these morphologies can be much more easily distinguished by other methods, such as the $\\mathcal{Q}$--parameter \\citep{Cartwright04}. Characterizing substructure in clusters is important, as this has recently been shown to facilitate dynamical mass segregation and disrupt binary and planetary systems in young clusters \\citep{Allison10,Parker11c,Parker12a}.  (ii) No single morphology is a good match to the observed $\\Sigma$ distribution presented in \\citet{Bressert10}. We interpret this as being due to the superposition of $\\Sigma$ values from the different star forming regions in the observational sample.  (iii) The $\\Sigma$ distribution can be significantly shifted to lower densities for both substructured and smooth clusters due to early dynamical evolution of the cluster. In the case of a supervirial (expanding) cluster, the median $\\Sigma$ value can be similar to the observed distribution after just 1\\,Myr of dynamical evolution. The dynamical evolution of a fractal causes the distribution of $\\Sigma$ to widen, as well as shifting it to lower densities. Even an initially dense cluster in virial equilibrium will expand, due to two-body relaxation and push the $\\Sigma$ distribution to lower values in the first 10\\,Myr of the cluster's life.  (iv) Substructure is quickly erased through dynamical interactions, and therefore the $\\mathcal{Q}$--parameter reaches high values rapidly. If all star forming regions start substructured, then a low $\\mathcal{Q}$--parameter could indicate that the $\\Sigma$ distribution reflects the primordial density of the region. This local substructure could facilitate dynamical processing of systems, without the need for high global densities.  (v) Although the dynamical scenarios that drastically alter the $\\Sigma$ distribution (a collapsing fractal cluster, or supervirial, initially dense Plummer sphere), can be viewed as extrema in terms of the formation and evolution of clusters, they represent possible initial conditions for the ONC \\citep{Kroupa99,Parker09a,Allison09b,Allison10,Parker11c}, which would contribute $\\sim$25\\,per cent of stars to the local volume-limited $\\Sigma$ distribution if it were to be included in the sample. If primordial substructure is taken into account, the fraction of stars that could be affected by dynamics in a cluster could be as high as 50\\,per cent.  (vi) If star forming regions form in virial equilibrium, then dynamical relaxation cannot significantly alter the $\\Sigma$ distribution. In such systems, the $\\Sigma$ distribution is a tracer of the initial density of a star forming region.  The combination of points (iii), (iv) and (v) lead us to suggest that the $\\Sigma$ distribution in local star forming regions may indicate a relatively quiescent natal environment for star formation, but the possibility of a more dynamically active initial environment for up to 50\\,per cent of the stars cannot be ruled out. We suggest that the $\\Sigma$ distribution in local star forming regions should not be over-interpreted \\citep[see also][]{Bate98b,Gieles12}, and that one should use a range of different metrics (e.g.\\,\\,a minimum spanning tree analysis such as the $\\mathcal{Q}$--parameter in tandem with age estimates, KS tests and``inverse binary population synthesis'' \\citep{Kroupa95a,Kroupa95b}) to determine the past and present state of star forming regions. Detailed kinematic information from the GAIA mission and its associated spectroscopic surveys will enable a detailed dynamical history of field and cluster stars to be made, which could in principle be used to infer the density of the birth environment of most stars."
    },
    "1208/1208.2330_arXiv.txt": {
        "abstract": "We discuss a novel sparsity prior for compressive imaging in the context of the theory of compressed sensing with coherent redundant dictionaries, based on the observation that natural images exhibit strong average sparsity over multiple coherent frames. We test our prior and the associated algorithm, based on an analysis reweighted $\\ell_1$ formulation, through extensive numerical simulations on natural images for spread spectrum and random Gaussian acquisition schemes. Our results show that average sparsity outperforms state-of-the-art priors that promote sparsity in a single orthonormal basis or redundant frame, or that promote gradient sparsity. Code and test data are available at https://github.com/basp-group/sopt. ",
        "introduction": "\\label{sec:intro} Compressed sensing (CS) introduces a signal acquisition framework that goes beyond the traditional Nyquist sampling paradigm~\\cite{fornasier11}. Consider a complex-valued signal $\\bm{x}\\in\\mathbb{C}^{N}$, assumed to be sparse in some orthonormal basis $\\mathsf{\\Psi}\\in\\mathbb{C}^{N\\times N}$, i.e.~$\\bm{x} = \\mathsf{\\Psi}\\bm{\\alpha}$ for $\\bm{\\alpha}\\in\\mathbb{C}^{N}$ sparse. Also consider the measurement model $\\bm{y}=\\mathsf{\\Phi}\\bm{x}+\\bm{n}$, where $\\bm{y}\\in\\mathbb{C}^{M}$ denotes the measurement vector, $\\mathsf{\\Phi}\\in\\mathbb{C}^{M\\times N}$ with $M<N$ is the sensing matrix, and $\\bm{n}\\in\\mathbb{C}^{M}$ represents noise. The most common approach to recover $\\bm{x}$ from $\\bm{y}$ is to solve the following convex problem \\cite{fornasier11}: $\\min_{\\bar{\\bm{\\alpha}}\\in\\mathbb{C}^{N}}\\|\\bar{\\bm{\\alpha}}\\|_{1} \\textnormal{ subject to }\\| \\bm{y}-\\mathsf{\\Phi \\Psi}\\bar{\\bm{\\alpha}}\\|_{2}\\leq\\epsilon$, where $\\epsilon$ is an upper bound on the $\\ell_{2}$ norm of the noise and $\\|\\cdot\\|_1$ denotes the $\\ell_{1}$ norm. The signal is recovered as $\\hat{\\bm{x}}=\\mathsf{\\Psi}\\hat{\\bm{\\alpha}}$, where $\\hat{\\bm{\\alpha}}$ denotes the solution to the above problem. Such problems, solving for the signal representation in a sparsity basis, are known as synthesis-based problems. Standard CS provides results if $\\mathsf{\\Phi}$ obeys a Restricted Isometry Property (RIP) and $\\mathsf{\\Psi}$ is orthonormal~\\cite{fornasier11}. However, signals often exhibit better sparsity in a redundant dictionary ~\\cite{gribonval03,bobin07,starck10}.  Recent works have begun to address CS with redundant dictionaries, i.e.~where $\\mathsf{\\Psi}\\in\\mathbb{C}^{N\\times D}$, with $N<D$, so that $\\bm{x} = \\mathsf{\\Psi}\\bm{\\alpha}$ with $\\bm{\\alpha}\\in\\mathbb{C}^{D}$. Rauhut et al.~\\cite{rauhut08} find conditions on $\\mathsf{\\Psi}$ such that $\\mathsf{\\Phi \\Psi}$ obeys the RIP to recover $\\bm{\\alpha}$ in a synthesis formulation. Cand\\`{e}s et al.~\\cite{candes10} provide a theoretical analysis of the $\\ell_1$ analysis-based problem. As opposed to synthesis, the analysis formulation solves for the signal itself: \\begin{equation}\\label{cs6} \\min_{\\bar{\\bm{x}}\\in\\mathbb{C}^{N}}\\|\\mathsf{\\Psi}^{\\dagger}\\bar{\\bm{x}}\\|_{1} \\textnormal{ subject to }\\| \\bm{y}-\\mathsf{\\Phi}\\bar{\\bm{x}}\\|_{2}\\leq\\epsilon, \\end{equation} where $\\mathsf{\\Psi}^{\\dagger}$ denotes the adjoint operator of $\\mathsf{\\Psi}$. The aforementioned work~\\cite{candes10} extends the standard CS theory to coherent and redundant dictionaries, providing theoretical stability guarantees based on a general condition of the sensing matrix $\\mathsf{\\Phi}$, coined the Dictionary Restricted Isometry Property (D-RIP). The D-RIP is a natural extension of the standard RIP. In fact many random matrices that obey the standard RIP also obey the D-RIP, like Gaussian or Bernoulli ensembles. Also, the subsampled Fourier matrix multiplied by a random sign matrix satisfies the D-RIP~\\cite{krahmer11}, which provides a fast sensing operator. Interestingly, this approach falls within the spread spectrum framework proposed in \\cite{puy12b}. If $\\mathsf{\\Phi}$ satisfies the D-RIP and $\\mathsf{\\Psi}$ is a general frame, Cand\\`{e}s et al.\\ prove in \\cite{candes10} that the solution to \\eqref{cs6}, denoted $\\hat{\\bm{x}}$, satisfies the following error bound: \\begin{equation}\\label{thm1} \\| \\hat{\\bm{x}}-\\bm{x}\\|_2\\leq C_0\\epsilon + C_1K^{-1/2}\\left \\| \\mathsf{\\Psi}^{\\dagger}\\bm{x}-(\\mathsf{\\Psi}^{\\dagger}\\bm{x})_K \\right\\|_1, \\end{equation} where $(\\mathsf{\\Psi}^{\\dagger}\\bm{x})_K$ denotes the best $K$-term approximation of $\\mathsf{\\Psi}^{\\dagger}\\bm{x}$ and $C_0$ and $C_1$ are numerical constants. Similar properties to the D-RIP coined $\\Omega$-RIP are introduced in \\cite{giryes13} in the context of the co-sparsity analysis model.  In \\cite{carrillo12} some of the authors of this paper proposed a novel sparsity analysis prior in the context of Fourier imaging in radio astronomy. Our approach relies on the observation that natural images are simultaneously sparse in various frames, in particular wavelet frames, or in their gradient, so that promoting average signal sparsity over multiple frames should be a powerful prior. In the present work, the average sparsity prior is put in the generic context of compressive imaging within the theory of CS with coherent redundant dictionaries. The associated reconstruction algorithm, based on an analysis reweighted $\\ell_1$ formulation, is dubbed Sparsity Averaging Reweighted Analysis (SARA). We evaluate SARA through extensive numerical simulations for spread spectrum and Gaussian acquisition schemes. Our results show that the average sparsity prior outperforms state-of-the-art priors. ",
        "conclusions": "\\label{sec:Conclusion} In this letter we have discussed the novel SARA regularization method and algorithm for compressive imaging in the theoretical context of CS with coherent redundant dictionaries. The approach relies on the observation that natural images exhibit strong average sparsity. We have evaluated SARA under two different acquisition schemes: spread spectrum and random Gaussian measurements. Experimental results demonstrate that the sparsity averaging prior embedded in the analysis reweighted $\\ell_1$ formulation of SARA outperforms state-of-the-art priors, based on single frame or gradient sparsity, both in terms of $\\mathrm{SNR}$ and visual quality. An MR imaging illustration also corroborates these conclusions for Fourier imaging. Code and test data are available at https://github.com/basp-group/sopt.  Future work will concentrate on finding a theoretical framework for the average sparsity model. Specialized results are indeed needed in the particular case of concatenation of frames for an estimate of the number of measurements required for accurate image reconstruction. It would be interesting to explore the connections between average sparsity and the co-sparsity model, which proposes a general framework for general analysis operators (see \\cite{giryes13} and references therein). Also, it was recently shown in \\cite{oymak13} that combinations of convex relaxation priors do not yield better results than exploiting only one of those priors, while non-convex approaches can exploit multiple models. Those results suggest that the re-weighting approach in SARA to approximate the non-convex $\\ell_0$ norm is fundamental to exploit average sparsity, as observed in the simulation results."
    },
    "1208/1208.5726_arXiv.txt": {
        "abstract": "The conversion of fast waves to the Alfv\\'en mode in a realistic sunspot atmosphere is studied through three-dimensional numerical simulations. An upward propagating fast acoustic wave is excited in the high-$\\beta$ region of the model. The new wave modes generated at the conversion layer are analyzed from the projections of the velocity and magnetic field in their characteristic directions, and the computation of their wave energy and fluxes. The analysis reveals that the maximum efficiency of the conversion to the slow mode is obtained for inclinations of 25 degrees and low azimuths, while the Alfv\\'en wave conversions peaks at high inclinations and azimuths between 50 and 120 degrees. Downward propagating Alfv\\'en waves appear at the regions of the sunspot where the orientation of the magnetic field is in the direction opposite to the wave propagation, since at these locations the Alfv\\'en wave couples better with the downgoing fast magnetic wave which are reflected due to the gradients of the Alfv\\'en speed. The simulations shows that the Alfv\\'en energy at the chromosphere is comparable to the acoustic energy of the slow mode, being even higher at high inclined magnetic fields. ",
        "introduction": "\\label{sect:introduction}  The study of solar oscillations has proven to be a powerful tool to infer the properties of the solar interior.  Global helioseismology, based on the interpretation of the eigenfrequencies of the resonant modes of oscillations, has provided a robust description of the internal structure and dynamics of the Sun. In the last years, these studies have been complemented by focusing on local features. Local helioseismology interprets the full wave field observed at the surface. Several complemetary techniques have been developed to probe local perturbations, as Fourier-Hankel spectral analysis \\citep{Braun1995}, ring-diagram analysis \\citep{Hill1988}, time-distance helioseismology \\citep{Duvall+etal1993}, and acoustic holography \\citep{Lindsey+Braun1990}.  For an accurate interpretation of the measurements obtained from all these procedures, it is extremely important to achieve a deep understanding of the physics involved in the wave propagation. Due to the increasing interest of helioseismologists in active regions, such as sunspots, the knowledge of the inteaction of waves with magnetic structures has been greatly developed in the recent years. The conversion from fast-mode high-$\\beta$ acoustic waves to slow-mode low-$\\beta$ waves in solar active region is well-understood from the theoretical point of view \\citep{Cally+Bogdan1993, Crouch+Cally2003, Schunker+Cally2006, Cally2006}. An example of the success of this analytical development is the comparison of the observational absorption and phase shift data \\citep{Braun+etal1988, Braun1995} with the modeled results obtained from the conversion \\citep{Cally+etal2003, Crouch+etal2005}. Numerically, the fast-to-slow conversion is also well studied in sunspot-like atmospheres \\citep{Bogdan+etal2003, Rosenthal+etal2002, Khomenko+Collados2006, Khomenko+etal2009, Felipe+etal2010a}, as well as flux tubes \\citep{Hasan+etal2003, Hasan+Ulmschneider2004}. See \\citet{Khomenko2009} for a review.  Not all the developments of phenomena associated with wave propagation in magnetic structures have reached the same degree of maturity. An insufficiently modelled effect is the fast-to-Alfv\\'en conversion. Pure Alfv\\'en waves can only exist in those mediums which are homogeneous in the direction perpendicular to both magnetic field and wavenumber. In general, this will not be the situation in a gravitationally stratifed atmosphere with complex magnetic field structure, although we will simply refer to Alfv\\'en waves even in this case. The fast-to-Alfv\\'en conversion only occurs when the wave propagation is not contained in the same plane of the density stratification and magnetic field and, thus, three dimensional (3D) analysis are necessary. The study of this process was started by \\citet{Crouch+Cally2005}, who studied the 3D propagation of oscillations in a polytrope permeated by an uniform magnetic field of arbitrary inclination, and found downward Alfv\\'en waves. \\citet{Cally+Goossens2008} obtained that the conversion to the Alfv\\'en mode is most efficient for field inclinations from vertical between 30 and 40 degrees, and azimuth angles (the angle between the magnetic field and wave propagation planes) between 60 and 80 degrees. \\citet{Cally+Hansen2011} found that the interaction between fast and Alfv\\'en waves is spread across many scale height, unlike the fast-to-slow, which is limited around the layer where the sound speed $c_S$ and the Alfv\\'en speed $v_A$ are similar. As the frequency increases the mode conversion is progressively more localized, although at the frequencis relevant to local helioseismology (around 3-5 mHz) the fast-to-Alfv\\'en conversion region spans the whole chromosphere.  Recently, Khomenko and Cally have studied the conversion to the Alfv\\'en mode by means of 2.5D numerical simulations in homogeneous field configurations \\citep{Khomenko+Cally2011}, as well as realistic sunspot-like structures \\citep{Khomenko+Cally2012}. In these works they obtained the dependence of the efficiency of the conversion with the inclination and azimuth. However, in the sunspot configuration the conversion was only evaluated at some limited angles corresponding to selected 2D planes of the model. The aim of this work is to extend the results from \\citet{Khomenko+Cally2012} to 3D. It will allow us to populate the voids in their diagrams of fast-to-Alfv\\'en conversion efficiency. More relevant, the development of full 3D simulations provides the results for the complete wave field, where waves can propagate freely in all spatial directions without being restricted to a two-dimensional plane. The conversion to the Alfv\\'en mode in realistic 3D atmospheres has not been studied before with the detailed evaluation of the efficiency discussed in this paper. ",
        "conclusions": "\\label{sect:scattering}  In this paper we have studied the conversion from fast to Alfv\\'en modes in a realistic sunspot-like atmosphere by means of 3D numerical simulations. This study provides a direct comparison with analytical works \\citep{Cally+Goossens2008, Cally+Hansen2011} as well as numerical simulations \\citep{Khomenko+Cally2011,Khomenko+Cally2012}.   With regards to the fast-to-slow conversion, the  efficiency of the transformation peaks at inclinations around 25$^o$ and low azimuths, and it decreases at higher azimuths. This result shows a good agreement with previous works \\citep{Schunker+Cally2006, Cally+Goossens2008} in homogeneous magnetic field configurations. However, the sunspot-like structure of this simulation produce some differences in the wave modes of the upper atmosphere. In the central and left part of the sunspot the simulation shows downward propagating slow waves, which are reflected due to the chromospheric increase of the temperature. A similar result was found in the 2.5D simulations from \\citet{Khomenko+Cally2012}, although in our case the downward flux near the axis is enhanced because of the higher increase of the temperature at the center of the sunspot. In \\citet{Khomenko+Cally2012} simulations the closest plane to the axis of the sunspot is located at a distance of 7.5 Mm, where the increase of the temperature in the transition region is lower (Figure \\ref{fig:temperature}).  The simulation presented in this paper supposes a step forward in the study of the conversion to the Alfv\\'en mode, since it accounts for full 3D. The work by \\citet{Khomenko+Cally2012} was limited to 2.5D, and only a few vertical slices across the sunspot model at several distances from the axis were analyzed. Since the fast-to-Alfv\\'en conversion depends strongly on the angle between the direction of wave propagation and the magnetic field, a full 3D simulation is required to retrieve the complete description. The simulation shows that the highest Alfv\\'enic wave energy is obtained at high inclinations and for azimuths between $50^o$ and $120^o$. The efficiency of the conversion obtained for the sunspot model shows a similar pattern to the homogeneous magnetic fields from \\citet{Cally+Goossens2008}, although the maximum of the magnetic energy of Alfv\\'en waves is shifted toward more inclined fields.  A comparison of the total energy in the upper atmosphere of the upward propagating slow acoustic energy and the corresponding to the Alfv\\'en mode reveals that the former one is around two times higher. This result remarks the important role of the Alfv\\'en energy at the chromosphere. At $\\theta\\approx 25^o$ the acoustic energy is much higher than the Alfv\\'en energy, especially at low azimuths where its conversion is more efficient. However, at higher inclinations the Alfv\\'en energy becomes dominant and at $\\theta=47^o$ and $\\phi$ around $90^o$ it is more than 7 times higher than the acoustic energy. The horizontal size of the computational domain prevent us to measure the energy at higher inclinations for these azimuths, but from the tendency shown in the computed region one would expect a strongest dominance of the Alfv\\'en energy at higher inclinations. Moreover, the vertical limitation of the domain also reduces the energy in the Alfv\\'en mode. At the 5 mHz frequency of this simulation the fast-to-Alfv\\'en conversion region spans over 20 scale heights \\citep{Cally+Hansen2011}. Since our computational box has a much lower size, only a small fraction of the conversion is completed in our domain. The same limitation affects the simulations from \\citet{Khomenko+Cally2012}. However, in our case we have found a higher energy in the Alfv\\'en wave due to the higher inclinations covered by our computational domain.  Despite of the dependence of the conversion with the azimuth, due to the stochastic direction of the wave propagation in the Sun, one would expect to find an efficient conversion to the Alfv\\'en wave at all the locations surrounding a sunspot with a certain inclination. Near the umbra-penumbra boundary, the waves whose attack angle is below $90^o$ will generate upward propagating Alfv\\'en waves, while the incident waves which propagate in the opposite direction will produce downward Alfv\\'en waves. In this way, in these regions upgoing and downgoing waves will coexist. On the other hand, at the higher inclinations of the penumbra, the energy flux of the Alfv\\'en mode is only composed by upward propagating waves for all the directions of incidence of the fast acoustic wave where the conversion to the Alfv\\'en wave is efficient, as shown by the right panel from Figure \\ref{fig:fluxXY} at $\\theta\\approx 45^o$. It differs from the situation at $\\theta\\approx 30^o$, where at some regions of the sunspot there is negative flux. In the case of solar observations, where the incident waves propagate in all directions, the Alfv\\'en waves in an annular region surrounding the sunspot with higher inclinations (around $\\theta\\approx 45^o$) will only consist on upward propagating waves.  The small time step imposed by the high chromospheric Alfv\\'en speed has been an important concern in the development of the numerical simulations presented in this paper. Some considerations have been taken into account in the numerical configuration regarding this issue. Firstly, a magnetic field strength of 900 G was adopted at the axis of the sunspot at the photosphere. This value is well below what might be expected in a mature sunspot. A more realistic field strength would lower the height of the $c_S=v_A$ and the fast mode reflection layers. This would allow the fast-to-Alfv\\'en conversion to be produced in a bigger region of the computational domain, and might generate an even higher Alfv\\'en energy flux at the top of the computational domain. Secondly, this top boundary has been set at the chromosphere. The recent work by \\citet{Hansen+Cally2012} has shown the interesting effects of the chromosphere-corona transition region on fast-to-Alfv\\'en conversion. They found that the reflection of the Alfv\\'en waves at the transition region \\citep{Uchida+Sakurai1975} is sensitive to the distance between the fast reflection point and the transition region. The Alfv\\'en flux that can reach the corona is increased when this distance is small. In the present work the computational box does not include the transition region, since it would produce an even smaller time step, compromising the realization of this simulation using a reasonable amount of computational time. It would be interesting to extend the simulations to the corona and adopt a realistic magnetic field strength, but it is out of the scope of this paper and shall be accounted for in forthcoming works."
    },
    "1208/1208.0429_arXiv.txt": {
        "abstract": "We analyzed correlations among the rest frame spectral peak energy $E_{\\rm p}$, the observed frame 64ms peak isotropic luminosity $L_{\\rm p}$ and the isotropic energy $E_{\\rm iso}$ for 13 Short Gamma Ray Burst (SGRB) candidates having the measured redshift $z$, $T_{\\rm 90}^{\\rm obs}/(1+z)<2$ sec and well determined spectral parameters. A SGRB candidate is regarded as a misguided SGRB if it is located in the 3-$\\sigma_{\\rm int}$ dispersion region from the best-fit function of the $E_{\\rm p}$--$E_{\\rm iso}$ correlation for Long GRBs (LGRBs) while the others are regarded as secure SGRBs possibly from compact star mergers. Using 8 secure SGRBs out of 13 SGRB candidates, we tested whether $E_{\\rm p}$--$E_{\\rm iso}$ and $E_{\\rm p}$--$L_{\\rm p}$  correlations exist for SGRBs. We found  that $E_{\\rm p}$--$E_{\\rm iso}$ correlation for SGRBs($E_{\\rm iso} =10^{51.42 \\pm 0.15}{\\rm erg~ s^{-1}} ({E_{\\rm p}}/{\\rm 774.5~keV})^{1.58 \\pm 0.28}$) seems to exist  with the correlation coefficeint $r=0.91$ and chance probability $p=1.5\\times10^{-3}$. We found also that  the $E_{\\rm p}$--$L_{\\rm p}$ correlation for SGRBs($L_{\\rm p} = 10^{52.29 \\pm 0.066}{\\rm erg~ s^{-1}} ({E_{\\rm p}}/{\\rm 774.5~keV})^{1.59 \\pm 0.11}$) is tighter than $E_{\\rm p}$--$E_{\\rm iso}$ correlation since $r=0.98$ and $p=1.5\\times10^{-5}$. Both correlations for SGRBs are dimmer than those of LGRBs for the same $E_{\\rm p}$ by factors $\\sim$100 ($E_{\\rm p}$--$E_{\\rm iso}$) and $\\sim$ 5($E_{\\rm p}$--$L_{\\rm p}$). Applying the tighter $E_{\\rm p}$--$L_{\\rm p}$ correlation for SGRBs to 71 bright BATSE SGRBs, we found that pseudo redshift $z$ ranges from $0.097$ to  $2.258$ with the mean $<z>$ of 1.05. The redshifts of SGRBs apparently  cluster at lower redshift than those of LGRBs ($<z>\\sim 2.2 $), which supports the merger scenario of SGRBs. ",
        "introduction": "\\label{sec:introduction}  For Long Gamma Ray Bursts (LGRBs), several observational correlations among the rest frame spectral peak energy $E_ {\\rm p}$, the peak isotropic luminosity $L_{\\rm p}$ and the isotropic energy $E_{\\rm iso}$ in the prompt emission phase have been proposed. $E_{\\rm p}$--$E_{\\rm iso}$ correlation was first reported by \\citet{Amati:2002} and argued by many authors \\citep{Sakamoto:2004,Lamb:2004,Amati:2006,Amati:2009,Yonetoku:2010}.  As for $L_{\\rm p}$, \\citet{Yonetoku:2004} reported a rather tight correlation between $E_{\\rm p}$ and the observed frame 1-second peak isotropic luminosity $L_{\\rm p}$. In 2004, the number of LGRBs with well determined redshifts and spectral parameters was only 16. Nevertheless, the correlation was found to be very tight: the linear correlation coefficient ($r$) between $\\log E_{\\rm p}$ and $\\log L_{\\rm p}$ is $0.958$ and the chance probability ($p$) is $5.3\\times 10^{-9}$. Several authors argued on the property of the $E_{\\rm p}$--$L_{\\rm p}$ correlation \\citep{Ghirlanda:2005a,Ghirlanda:2005b, Krimm:2009} and confirmed the existence. \\citet{Tsutsui:2009b} found that adding a new observables $T_ {\\rm L}$, the luminosity time defined by $T_{\\rm L}=E_{\\rm iso}/L_{\\rm p}$, improves the correlation and gave $E_{\\rm p}$--$T_{\\rm L}$--$L_{\\rm p}$ correlation. In $E_{\\rm p}$--$T_{\\rm L}$--$L_{\\rm p}$ correlation, the intrinsic dispersion is reduced by $\\sim$ 40 \\% compared with the $E_{\\rm p}$--$E_{\\rm iso}$ and $E_{\\rm p}$--$L_{\\rm p}$ correlations.  \\citet{Ghirlanda:2004} applied the $E_{\\rm p}$--$L_{\\rm p}$ correlation to bright short Gamma Ray Bursts (SGRBs) observed by BATSE without measured redshift. That is, they assumed that SGRBs obey the same $E_{\\rm p}$--$L_{\\rm p}$ correlation of LGRBs and estimated the pseudo redshifts of SGRBs although no evidence for the existence of the $E_{\\rm p}$--$L_{\\rm p}$ correlation for SGRBs at that time. They found that the pseudo redshifts are obtained for all selected SGRBs and the distribution is similar to that of LGRBs known at that time. On the other hand, \\citet{Nakar:2005,Band:2005,Butler:2007,Shahmoradi:2010} argued that $E_{\\rm p}$--$L_{\\rm p}$ correlation might be due to selection effects, since $E_{\\rm p}$ was determined from the time integrated spectra. However, \\citet{Ghirlanda:2010} showed that in the individual pulses of several LGRBs, $E_{\\rm p}$--$L_{\\rm p}$ correlation holds for each pulse even though $E_ {\\rm p}$ changes an order of magnitude from pulse to pulse. Similar property was found for GRB061007 by \\citet{Ohno:2009}. These results strongly suggest that $E_{\\rm p}$--$L_{\\rm p}$ correlation is  not a result of selection biases but a real physical one.  As for SGRBs, the number of SGRBs with measured redshifts and $E_{\\rm p}$ was so small that it was difficult to check if $E_{\\rm p}$--$L_{\\rm p}$ correlation holds or not. However, \\citet{Ghirlanda:2011} showed that for 14 Fermi/GBM SGRBs without redshifts, the individual pulses follow a relation of $E_{\\rm p}$ $\\propto F_{\\rm pulse}^s$ with $s\\sim 1$ where $F_{\\rm pulse}$ is the observed energy flux. This reminds us what happened to the individual pulses of LGRBs in \\citet{Ghirlanda:2010} and suggests that a similar correlation might exist even for SGRBs in the rest frame.  In this study, we select 13 SGRB candidates with well determined redshift,  spectral parameters, $E_ {\\rm p}$, $L_ {\\rm p}$ and $E_{\\rm iso}$ to see if the correlations among $E_ {\\rm p}$, $L_ {\\rm p}$ and $E_{\\rm iso}$ exist. In section 2, we will show that our criteria on SGRBs yield 8 secure SGRBs out of 13 SGRB candidates. Using these SGRBs, we examine if the $E_{\\rm p}$--$E_{\\rm iso}$ and $E_{\\rm p}$--$L_{\\rm p}$ correlations exist or not. In section 3, we will apply the $E_{\\rm p}$--$L_{\\rm p}$ correlation obtained in section 2 to 71 bright BATSE SGRBs without measured redshift to determine the pseudo redshift $z$. Section 4 will be devoted to discussions. Throughout the paper we adopt a cosmological model with $\\Omega_\\Lambda=0.7 $, $\\Omega_m=0.3 $ and $H_0=70{\\rm km s^{-1}Mpc^{-1}}$ ",
        "conclusions": "As mentioned before, the comparison of the distribution of LGRBs and SGRBs in $E_{\\rm p}$--$E_{\\rm iso}$ and $E_{\\rm p}$--$L_{\\rm p}$ planes was performed in \\citet{Zhang:2012}. As to the former, they recognized the difference in the distribution and found that the $E_{\\rm p}$--$E_{\\rm iso}$ correlation from SGRBs is almost parallel but dimmer by a factor of 10 compared with the one from LGRBs. This is reasonably consistent with our result. However, as to the latter, they insisted that SGRBs follow the same correlation as the one derived from LGRBs, which is in contradiction with our analysis. Below, we will discuss the origin of this discrepancy.  The left panel of Figure \\ref{vszhang12} shows \\yonetoku diagram for our secure SGRB sample (red filled circles) and SGRB sample from \\citet{Zhang:2012} (blue filled squares) with the best-fit line for each sample. Here it should be noted that the best-fit line solely from SGRBs was not derived in \\citet{Zhang:2012} and was newly derived here. The best-fit line for LGRBs of \\citet{Yonetoku:2010} is also plotted with a black solid line. We can see that all but one events are located below the LGRB line. This fact indicates that SGRBs are systematically dimmer than LGRBs with the same $E_{\\rm p}$ even if we consider SGRB sample by \\citet{Zhang:2012}.   Here we should note that 7 of 8 SGRBs in our sample are actually the same events with \\citet{Zhang:2012}, though they have different values of $L_{\\rm p}$ which leads to the different best-fit lines. Let us comment on the difference in each event. First of all, our peak luminosities are uniformly defined by the 64-msec time resolution in obesrver frame for all SGRBs, while \\citet{Zhang:2012} used different time resolutions (from 4 ms to 1024 ms). This is because they adopted the values of $L_{\\rm p}$ reported by multiple observation teams. The value of 4-msec peak luminosity is typically a few times larger than that of 1024-msec peak luminosity \\citep{Tsutsui:2011,Tsutsui:2012a} so that using different time resolution to define $L_{\\rm p}$ would make artificial dispersion in the $E_{\\rm p}$--$L_{\\rm p}$ correlation \\citep{Yonetoku:2010}. Secondly, we integrate energy spectra between 1-100,000 keV in GRB rest frame to calculate $L_{\\rm p}$, while an energy range of 1-10,000 keV was considered in \\citet{Zhang:2012}. Therefore they tend to underestimate $L_{\\rm p}$ compared with us. From these reasons, the values of $L_{\\rm p}$ are different in the two samples and, we believe, our sample is more reliable compared with \\citet{Zhang:2012}.   On the other hand, the correlation for the LGRBs are also different between our analysis and \\citet{Zhang:2012}. The right panel of Figure \\ref{vszhang12} is the same diagram as the left panel, but the best-fit line is for LGRBs in \\citet{Ghirlanda:2010} which \\citet{Zhang:2012} uses. The best-fit line for combined short and long GRB sample obtained by \\citet{Zhang:2012} is also plotted with a blue dash-doted line. Although the best-fit lines are significantly different from ours, the same tendency can still be seen.   It is beyond the scope of this paper to explain the difference between the \\yonetoku correlations for LGRBs from \\citet{Ghirlanda:2010} and the one from \\citet{Yonetoku:2010}, and we just make a short remark on this. The major difference comes from the treatment of GRB 060218. The former regarded it as a ordinary LGRB, while in the latter it was regarded as an outlier by a statistical argument. Because GRB 060218 is located far away from the \\yonetoku correlation of \\citet{Yonetoku:2010} (more than 8-$\\sigma$), it makes the best-fit line much steeper like the one of \\citet{Ghirlanda:2010}. Anyway, it seems to be robust that SGRBs have systimatically smaller $L_{\\rm p}$ than LGRBs for a given $E_{\\rm p}$, even if we consider the possible systematic errors in LGRBs, as well as SGRBs.   \\begin{figure*} \\rotatebox{0}{\\includegraphics[width=75mm]{fig5-1.eps}} \\rotatebox{0}{\\includegraphics[width=75mm]{fig5-2.eps}} \\caption{(Left) The $E_{\\rm p}$--$L_{\\rm p}$ diagram for SGRBs. Our eight secure SGRBs are marked with red filled squares and seven SGRBs from Zhang et al. (2012) with blue filled triangles. The best fit for each sample are indicated with red solid line and blue solid line, respectively and the black line is the best fit for LGRBs from \\citet{Yonetoku:2010}. (Right) The same diagram as the left panel, but the black line is the the best fit for LGRBs from \\citet{Ghirlanda:2010} and blue dash-doted line is the one for combined LGRBs and SGRBs from \\citet{Zhang:2012}. } \\label{vszhang12} \\end{figure*}  In this paper, we suggested possible correlations among $E_{\\rm p}$, $L_{\\rm p}$ and $E_{\\rm iso}$  even for SGRBs. However, the correlations for SGRBs are much dimmer than those for LGRBs. The $E_{\\rm p}$--$E_{\\rm iso}$ ($E_{\\rm p}$--$L_{\\rm p}$) correlation for SGRBs is located approximately $10^2$ ($5$) times below the one for LGRBs. For the $E_{\\rm p}$--$L_{\\rm p}$ correlation for SGRBs, similar arguments have been made by some authors \\citep{Ghirlanda:2009,Zhang:2012}, but we for the first time argue that there exist distinct $E_{\\rm p}$--$E_{\\rm iso}$ and $E_{\\rm p}$--$L_{\\rm p}$ correlations for SGRBs.  The distinction between SGRBs and LGRBs becomes much clearer if we use the gold sample of LGRBs compiled by \\citet{Tsutsui:2012a}. \\citet{Tsutsui:2012a} argued that there are two $E_{\\rm p}$--$L_{\\rm p}$ correlations, one is for small-$ADCL$  GRBs and the other is for large-$ADCL$ GRBs, where $ADCL$ stands for Absolute Deviations from  Constant Luminosity. In figure \\ref{Ep-Lp}, we shows the $E_{\\rm p}$--$L_{\\rm p}$ diagram for small-$ADCL$ LGRBs (black filled circles), large-$ADCL$ LGRBs (blue filled triangles), and secure SGRBs (red filled squares). The outliers of gold sample in \\citet{Tsutsui:2012a} and misguided SGRBs are removed from this figure. We can see the existence of three distinct $E_{\\rm p}$--$L_{\\rm p}$ correlations with almost the same power law index and different amplitudes.  \\begin{figure} \\rotatebox{0}{\\includegraphics[width=85mm]{fig6.eps}} \\caption{The $E_{\\rm p}$ -- $L_{\\rm p}$ diagram both for short and long GRBs. The LGRBs from gold sample in \\citet{Tsutsui:2012a} are marked with black filled circles (small-$ADCL$ GRBs) and blue filled triangles (large-$ADCL$ GRBs). The outliers of gold sample in \\citet{Tsutsui:2012a} and misguided GRBs are removed from this figure. } \\label{Ep-Lp} \\end{figure}  The accurate functional forms of $E_{\\rm p}$--$E_{\\rm iso}$ and $E_{\\rm p}$--$L_{\\rm p}$ correlation are very important to study the progenitor and the radiation mechanism of SGRBs. At present the intrinsic dispersion is rather large, that is , 0.13(0.39) in logarithm for $E_{\\rm p}$--$L_{\\rm p}$($E_{\\rm p}$--$E_{\\rm iso}$), respectively. This is mainly due to the small number of secure SGRBs, which prevents more detailed analysis. In conclusion we need more data of SGRBs with accurate $z$, $E_ {\\rm p}$, $L_ {\\rm p}$ and $E_{\\rm iso}$ to confirm or refute the $E_{\\rm p}$--$E_{\\rm iso}$ and $E_{\\rm p}$--$L_{\\rm p}$ correlations for SGRBs suggested in this Letter."
    },
    "1208/1208.3931_arXiv.txt": {
        "abstract": "We obtain approximate analytic expressions for the critical value of the total angular momentum of a non-relativistic test particle moving in the Kerr geometry, such that it will be captured by the black hole.   The expressions apply to arbitrary orbital inclinations, and are accurate over the entire range of angular momentum for the Kerr black hole.   The expressions can be easily implemented in N-body simulations of the evolution of star clusters around massive galactic black holes, where such captures play an important role. ",
        "introduction": "\\label{sec:intro}   The behavior of large numbers of stars orbiting around a central massive rotating black hole is a subject of great current interest (for recent reviews, see \\cite{alexander05,genzel10}).   Recently, Merritt {\\em et al.}\\cite{mamw1}  used N-body simulations to construct and evolve a cluster of stars around a massive black hole such as the one at the center of our own galaxy, and studied the perturbations induced by those stars on the orbits of hypothetical stars that might be used to test the black hole no-hair theorems \\cite{cmwnohair,lalehcliff}.  In separate work, Merritt {\\em et al.} also  analyzed the rate of injection of stars onto highly eccentric orbits around a massive black hole \\cite{mamw2}.  Such orbits would lead to an ``extreme mass-ratio inspiral\" (EMRI) that would result in the emission of detectable gravitational radiation.  Generally, such N-body simulations must be evolved for long periods of time in order to allow for all correlations and relaxation mechanisms to be established, so as to develop realistic estimates for rates of particular events, such as EMRI production. In the case of \\cite{mamw2}, the timescale was 10 Myr.  The simulations of Merritt {\\em et al.}~\\cite{mamw1} included the Newtonian gravitational attraction between the stars and the black hole and between each other, as well as post-Newtonian effects associated with the massive, spinning central black hole, including the dragging of inertial frames and the quadrupole moment of the black hole's field.  The simulations of \\cite{mamw2} restricted attention to a non rotating black hole, but included the capture of stars by the black hole.  A star-star encounter can place one star on an orbit with sufficiently small angular momentum that it can be captured directly by the black hole.   This is a non-trivial effect: in typical simulations carried out in \\cite{mamw2}, out of 50 stars in the cluster, 17 were captured by $t = 2$ Myr, and 30 were captured by $t= 10$ Myr.  In \\cite{mamw2} the condition for a capture was simple: if the total angular momentum per unit mass $L$ of a star after an encounter with another star was smaller than the critical value $L_c = 4GM/c$, where $M$ is the mass of the black hole, the star was deemed to be captured, the mass of the black hole was augmented by the star's mass, and the hole was given a suitable recoil.   This value of $L_c$ corresponds to the well-known angular momentum of an unstable circular orbit in the Schwarzschild geometry with conserved relativistic energy per unit mass $E =1$.   Grossman {\\em et al.}\\cite{levin2012} call this the  innermost bound spherical orbit (IBSO). The latter condition is appropriate here, because the stars are very far from the black hole in non-relativistic orbits when they encounter one another.  Accordingly, one can treat them as being barely bound ($E=1$) from the point of view of the black hole's strong gravity.  However, there is every reason to expect that the typical massive black hole at the center of a galaxy, including our own, will be rotating, perhaps even very rapidly rotating.  As a consequence, the capture of a star will not be isotropic: stars in prograde orbits will be less likely to be captured than stars in retrograde orbits.   For example, for $E=1$ stars with orbits restricted to the equatorial plane, the critical angular momentum per unit mass for capture is given by \\begin{equation} L_c = \\frac{GM}{c} \\left ( 2 + 2 \\sqrt{1\\mp \\tilde{a}} \\right ) \\,, \\end{equation} where $\\tilde{a}$ is the dimensionless Kerr parameter, related to the black hole's angular momentum $J$ by $\\tilde{a} \\equiv Jc/GM^2 \\, (0 \\le \\tilde{a} \\le 1)$. The upper(lower) sign corresponds to prograde(retrograde) orbits (see, for example, \\cite{bardeen73,young76}).   This anisotropic capture could have important consequences for the evolution of the cluster, imparting a net angular momentum to it, for instance, because of the preferential loss of stars moving in a retrograde manner relative to the black hole.  For orbits out of the equatorial plane, the problem is much more complicated, despite the presence of the additional  ``Carter'' constant of motion, which makes the mathematical problem completely integrable.  Generally one must resort to numerical solutions to study the behavior of orbits near this critical point \\cite{Glampedakis,levin}.  Instead, we have found an analytic expression for the critical angular momentum for arbitrary orbits that is approximate, but surprisingly accurate over almost the whole range of black hole spins.  It has the form \\begin{equation} L_c = \\frac{GM}{c} \\left ( 2 + 2 \\sqrt{1 - \\tilde{a} \\cos i - \\frac{1}{8}\\tilde{a}^2 \\sin^2 i F(\\tilde{a}, \\cos i) } \\right ) \\,. \\label{Lcrit0} \\end{equation} Because the total angular momentum is no longer a conserved quantity, we {\\em define} $L$ in terms of the Carter constant: $L \\equiv \\sqrt{C}$ (all quantities are suitably scaled so as to be independent of the mass of the star).  We also define $\\cos i \\equiv L_z/L$, where $L_z$ is the conserved $z$-component of angular momentum per unit mass.  In the Schwarzschild limit and the Newtonian limit,  $L$ {\\em is} the total conserved angular momentum and $i$ is the orbital inclination, with $0 \\le i \\le \\pi/2$ corresponding to prograde orbits ($L_z >0$) and $\\pi/2 \\le i \\le \\pi$ corresponding to retrograde orbits ($L_z < 0$).  The function $F(\\tilde{a}, \\cos i) = 1 + (\\tilde{a}/2) \\cos i + \\dots  $  is a series expansion in $\\tilde{a}$ shown below in Eq.\\ (\\ref{Fseries}).   With the series carried to $O(\\tilde{a}^4)$, the solution (\\ref{Lcrit0}) agrees with numerical solutions for the critical angular momentum to better than $0.5$ percent for $0 \\le \\tilde{a} \\le 0.9$; and to better than $5$ percent for $0.9 \\le \\tilde{a} \\le 0.99$.  For the special case of the extreme Kerr black hole ($\\tilde{a} =1$), we obtain a separate, but very accurate analytic fit for $L_c$ as a function of $\\cos i$.    Equation (\\ref{Lcrit0}) can be easily implemented in an N-body code as a capture criterion.  The rest of this paper provides details.  In Sec.\\ \\ref{basic} we review the basic equations for motion in the Kerr geometry, and in Sec.\\ \\ref{critical} we obtain the critical value of $L$ for capture in the case of $E=1$.  Concluding remarks are made in Sec.\\ \\ref{conclude}.  Henceforth, we use units in which $G=c=1$. ",
        "conclusions": "\\label{conclude}  Implementation of this capture condition in an $N$-body code will depend on where along a given orbit the star's orbit elements are determined in order to find $r_p$, $i$ and $L_z$.  Presumably it is after the last important stellar encounter.  If the star is still relatively far from the black hole, then one can argue that relativistic corrections will be small compared to the variations shown in Eq.\\ (\\ref{Lcrit}).  If the star is close to the black hole when this evaluation is made, then it may be necessary to be more careful in applying Eq.\\ (\\ref{Lcrit}), for example by incorporating post-Newtonian corrections in the evaluation of the constants of motion for the orbit.  Given values of position and velocity of the star following the encounter, this can be done in a straightforward manner.  It would be interesting to try to generalize the capture condition to value of $E$ less than unity, so as to handle more relativistic clusters.   \\ack This work was supported in part by the National Science Foundation, Grant Nos.\\ PHY 09--65133 \\& 12--60995.  We thank the Institut d'Astrophysique de Paris for its hospitality during the completion of this work.  We are grateful to Leor Barack, Emanuele Berti, Kostas Glampedakis, Shahar Hod, Scott Hughes, Janna Levin, and David Merritt for useful comments."
    },
    "1208/1208.1261_arXiv.txt": {
        "abstract": "We study the collimation of relativistic magnetohydrodynamic jets by the pressure of an ambient medium, in the limit where the jet interior loses causal contact with its surroundings. This follows up a hydrodynamic study in a previous paper, adding the effects of a toroidal magnetic field threading the jet. As the ultrarelativistic jet encounters an ambient medium with a pressure profile with a radial scaling of $p~\\propto~r^{-\\eta}$ where $2<\\eta<4$, it loses causal contact with its surroundings and forms a boundary layer with a large pressure gradient. By constructing self-similar solutions to the fluid equations within this boundary layer, we examine the structure of this layer as a function of the external pressure profile. We show that the boundary layer always becomes magnetically dominated far from the source, and that in the magnetic limit, physical self-similar solutions are admitted in which the total pressure within the layer decreases linearly with distance from the contact discontinuity inward. These solutions suggest a `hollow cone' behavior of the jet, with the boundary layer thickness prescribed by the value of $\\eta$. In contrast to the hydrodynamical case, however, the boundary layer contains an asymptotically vanishing fraction of the jet energy flux. ",
        "introduction": "The outflows from active galactic nuclei (AGN) are thought to be highly relativistic (\\citealp{MB84}) and highly collimated (e.g. \\citealp{Jorstad05}), but the cause of this collimation is uncertain.  Because jet-launching is generally believed to be electromagnetically driven (e.g. \\citealp{Blandford82,Contopoulos94}), one of the most commonly-accepted explanations for the observed collimation is that jets are threaded with magnetic fields that cause collimation via magnetic tension (e.g. \\citealp{Benford78,MB95}). Supporting this theory, it has been demonstrated that both relativistic and non-relativistic hydromagnetic outflows must eventually become collimated (\\citealp{Chiueh91,Heyvaerts89}). For magnetic fields acting alone, however, collimation will only happen on extremely large scales (\\citealp{Eichler93,MB94,MB95}).  To cause jets to collimate on reasonable scales, there must be an additional mechanism at work. A logical culprit is confinement by the pressure of an external medium. Pressure confinement has been demonstrated to act effectively on its own (e.g. \\citealp{Levinson00,BL07,Kohler12}), and accretion disk winds surrounding an AGN provide an ideal ambient medium to help to collimate the jet.  There have been many numerical studies of magnetized jets (e.g. \\citealp{KomissarovNumSim99,Hawley06,Beckwith08a,McKinney09}), with the goal of forming a self-consistent description of the jet-launching and collimation mechanisms. These numerical simulations have several restrictions, however, one of which being that the boundary of the jet, rather than having its shape determined by pressure balance, is generally treated as a rigid wall (e.g. \\citealp{KomBarkVla07,KomVlaKon09,Komissarov11,Tchekhovskoy10}). This construction doesn't allow the ambient pressure to affect collimation of the jet.  Treatments that do include effects of the external medium commonly focus on describing jets that remain in causal contact (e.g. \\citealp{Zakamska08,Lyubarsky11}). As an ultrarelativistic jet expands into an ambient medium with a pressure profile $p \\propto r^{-\\eta}$, it will eventually lose causal contact if $\\eta > 2$ and the opening angle is greater than $1/\\Gamma$, where $\\Gamma$ is the bulk Lorentz factor of the fluid. Observations of gamma-ray bursts indicate that these relativistic jets largely have opening angles greater than $1 / \\Gamma$ (e.g. \\citealp{Piran04}; see \\citealp{Tchekhovskoy10} for discussion), and AGN outflows with large Lorentz factors may similarly be causally disconnected; thus the poorly-studied regime of a jet that has lost causal contact is of physical interest.  In a previous paper (\\citealp{Kohler12}, hereafter KBB12) we developed a model describing the recollimation boundary layer of a purely hydrodynamic, ``hot'' (pressure-dominated) jet with an ultrarelativistic equation of state. In this model, we assumed that the pressure outside the jet decreases with $r$ so rapidly that the jet interior loses causal contact with its boundary, resulting in a shocked boundary layer forming within the jet. Though the jet interior is causally disconnected, the boundary layer is nevertheless narrow enough to remain in causal contact itself. Assuming self-similarity as a function of $r$, we calculated how the transverse structure of the jet boundary layer depends on the value of $\\eta$ in the external pressure profile.  We now expand this work to include, in addition to collimation by the external medium, the effects of a magnetic field within the jet. We include only a toroidal field, as it is the toroidal field that dominates the dynamics at large radii, far outside the light cylinder (\\citealp{MB84,Contopoulos95,Beskin09}). In Section 2, we first demonstrate that seeding a jet with a magnetic field at the base will always cause it to become magnetically dominated at large radii. We then find self-similar solutions for the boundary layer of the jet in the limit of magnetic dominance. In Section 3 we discuss the results, and in Section 4 we conclude. ",
        "conclusions": "We have evaluated the structure of a boundary layer within a hot, magnetohydrodynamic jet with an ultrarelativistic equation of state. We assumed that the jet as a whole is causally disconnected from its surroundings, but the boundary layer is thin and therefore in causal contact. We examined the impact on jet collimation of a toroidal magnetic field within the jet as well as an ambient medium with a pressure profile of $p \\propto r^{-\\eta}$, with $2<\\eta<4$.  We first demonstrated that the basic RMHD equations can be used to show that any jet boundary layer seeded with a toroidal magnetic field at its base will eventually become magnetically dominated at large radii.  We then constructed self-similar solutions for the boundary layer in the limit where the jet pressure is dominated by magnetic pressure. We found a special case of physical solutions where the jet pressure decreases linearly across the boundary layer, dropping to zero at a location set by the boundary conditions and the value of the pressure profile parameter $\\eta$. The boundary layer thickness is dependent upon the value of $\\eta$, with increasing $\\eta$ producing an increasingly wide boundary layer.  We further found that the thickness of the boundary layer decreases with radial distance, and the boundary layer contains a decreasing amount of the jet energy. This suggests that the addition of a magnetic field fundamentally changes the jet at large radius: whereas in the hydrodynamic case the inner boundary of the layer is a shock front through which material enters the layer, in the magnetohydrodynamic case the layer is bounded on the inside by a rarefaction front through which material leaves the boundary layer and rejoins the interior of the jet.  We found the position of this rarefaction front to occur just inside the layer from the location where the pressure vanishes, providing a smooth transition between the boundary layer and the jet interior and allowing for matching across the rarefaction front to the conditions in the jet interior. This matching would prescribe the values of the boundary conditions within the layer, and could potentially yield a solution where the rarefaction front is gradually decollimating, intercepting fewer and fewer streamlines as radial distance from the source increases.  In spite of the thinning of the boundary layer with radial distance, a sharp pressure gradient is nonetheless maintained across the layer, causing it to function as an insulating buffer between the jet interior and the ambient medium. Unlike the hydrodynamic case, the solutions for the structure of a magnetized jet do not have clear observational implications. Though the boundary layer contains a decreasing amount of energy as one looks further from the source, the layer might nonetheless have a high emissivity, which could be observationally important if the flow within the boundary layer is pointed along our line of sight.  The results presented in this paper provide the premise for a more complete treatment in numerical simulations of the effects of the ambient medium on collimation, both by demonstrating the behavior of the jet when the outer wall is allowed to change its shape, and by providing models that can be used to assess the effects of numerical resolution on simulation outcomes.  Ultimately, these results provide a foundation for future work examining energy dissipation in magnetized jets and the associated radiative observational signatures."
    },
    "1208/1208.4358_arXiv.txt": {
        "abstract": "{Submillimeter spectroscopic observations of comets provide an important tool for understanding their chemical composition and enable a taxonomic classification.} {We aim to determine the production rates of several parent- and product volatiles and the \\ce{^{12}C}/\\ce{^{13}C} isotopic carbon ratio in the long-period comet \\machholz{}, which is likely to originate from the Oort Cloud.} {The  line emission from several molecules in the coma was measured with high signal-to-noise ratio in January 2005 at heliocentric distance of 1.2 AU by means of high-resolution spectroscopic observations using the Submillimeter Telescope (SMT) at the Arizona Radio Observatory (ARO).} {We have obtained production rates of several volatiles (\\detected{}) by comparing the observed and simulated line-integrated intensities. We calculated the synthetic profiles using a radiative transfer code that includes collisions between neutrals and electrons, and the effects of radiative pumping of the fundamental vibrational levels by solar infrared radiation.  Furthermore, multiline observations of the \\ce{CH3OH} $J$ =  7--6 series allow us to estimate the rotational temperature using the rotation diagram technique.  We find that the \\ce{CH3OH} population distribution of the levels sampled by these lines can be described by a rotational temperature of \\rottemp{}. Derived mixing ratios relative to hydrogen cyanide are CO/\\ce{CH3OH}/\\ce{H2CO}/CS/HNC/\\ce{H^{13}CN}/HCN = 30.9/24.6/4.8/0.57/0.031/0.013/1 assuming a pointing offset of 8\\arcsec\\ due to the uncertain ephemeris at the time of the observations and the telescope pointing error.} {The measured relative molecular abundances in \\machholz\\ are between low- to typical values of those obtained in Oort Cloud comets, suggesting that it has visited the inner solar system previously and undergone thermal processing.  The HNC/HCN abundance ratio of $\\sim 3.1$\\% is comparable to that found in other comets, accounting for the dependence on the heliocentric distance, and could possibly be explained by ion-molecule chemical processes in the low-temperature atmosphere.  From a tentative \\ce{H^{13}CN} detection, the measured value of \\cratio\\ for the \\ce{H^{12}CN}/\\ce{H^{13}CN} isotopologue pair is consistent with a telluric value.  The outgassing variability observed in the HCN production rates over a period of two hours is consistent with the rotation of the nucleus derived using different observational techniques. } ",
        "introduction": "\\label{sec:intro}  Comets spend most of their lifetime in the outer solar system and therefore have not undergone much thermal processing.  Line emission from cometary atmospheres at submillimeter and radio wavelengths is a very useful tool for studying their physical and chemical conditions and relation with other bodies in the solar system \\citep{2002EM&P...90..323B,2004come.book..391B}. The coma structure and expansion velocity can be derived by fitting the observed line shapes using a molecular excitation code.  In addition, mixing ratios of volatiles such as \\ce{CH3OH}, CO and CS can be compared with observed chemical abundances in protoplanetary disks to improve our understanding of planet formation processes.  The composition of comets has been investigated in the last two decades to develop a classification based on abundances of primary chemical species that displays a great compositional diversity \\citep{1995Icar..118..223A,2003AdSpR..31.2563M,2004come.book..391B}. More than 20 parent volatile species that release directly from ices in the nucleus, in addition to several radicals and ions formed by photodissociation in the coma, have been detected via ground-based spectroscopic surveys at infrared and submillimeter wavelengths and in situ measurements.  The composition of some cometary ices show strong evidence of processing in the solar nebula and can provide clues about their place of formation and subsequent evolution.  The abundance of HCN relative to water has been observed to be roughly constant with a value of 0.1\\% in several comets for a wide range of heliocentric distances \\citep{2002EM&P...90..323B}.  Other species show a wide spread of production rates. For instance, \\ce{CH3OH} has been found to have a variable abundance relative to water, ranging between less than 0.15\\% to 6\\% at different heliocentric distances \\citep{2004come.book..391B}. There is no evident correlation between the observed relative abundances and the dynamical class of the comets.  Water is the most dominant volatile species and is typically used to determine relative abundances. Although water is not directly accessible from the ground at submillimeter wavelengths, it has been observed from space using the Submillimeter Wave Astronomical Satellite (SWAS), Odin and \\herschel\\ satellites \\citep{2000ApJ...539L.151N,2003A&A...402L..55L,2009P&SS...57.1596H} or inferred by observations of the hydroxyl (OH) radical at radio wavelengths \\citep{2002A&A...393.1053C}.  \\herschel{}'s Heterodyne Instrument for the Far Infrared (HIFI) is able to determine water production rates accurately \\citep{2010A&A...518L.150H}. Direct measurements of water can also be performed using ground-based telescopes by observations of non-resonance fluorescence emission at infrared wavelengths \\citep{1986Sci...232.1523M,2009ApJ...699.1563B}.   Comet \\object{\\machholz{}} was discovered on 27 August 2004 by Donald~E.~Machholz \\citep{2004IAUC.8394....1M}.  The comet passed perihelion on 24 January 2005 at a heliocentric distance of $\\rh = 1.205$~AU and geocentric distance of $\\Delta = 0.435$~AU.  During its closest approach to Earth at a distance of 0.35 AU on 6 January 2005 it reached a naked-eye visual magnitude of $m_\\mathrm{v}\\sim3.5$ as reported in the International Comet Quarterly.  Owing to its favorable viewing geometry in the northern hemisphere, the comet was extensively observed from the ground at various wavelengths.  \\machholz{} is a long-period highly eccentric comet whose origin is most likely the Oort Cloud according to the comet classification scheme by \\citet{1996ASPC..107..173L}, with an approximate initial orbital period of 118\\,000~years and eccentricity of 0.9994658 before the comet was perturbed gravitationally in the inner solar system \\citep{2004U31, NK1352}.  These values are close to the perihelion osculating elements \\citep[see][JPL Small-Body Database\\footnote{\\url{http://ssd.jpl.nasa.gov/sbdb.cgi?sstr=C/2004+Q2}}]{2005MPC..54558...1M}. Its dynamical classification is still a matter of debate because of the strong non-gravitational forces that make a backward orbital integration very unreliable.   In this paper we present high-resolution spectroscopic observations of several volatiles from comet \\machholz{} acquired at the Submillimeter Telescope (SMT).  Seven species are detected, namely \\sure{}, and a marginal detection of \\ce{H^{13}CN}. The comet was observed shortly pre-perihelion in January 2005 when it was at a distance $\\sim 0.36$~AU from Earth.  These observations provide information about the outgassing of several molecules and an isotopologue of HCN relative to HCN, which is often used as a proxy for water in cometary taxonomies.  We calculate the \\ce{CH3OH} rotational temperature of the ground vibrational level from several rotational lines and production rates for the observed molecules using a radiative transfer code to fit the observed line intensities.  Section~\\ref{sec:observations} presents our SMT observations of comet \\machholz{} and the reduction method.  In Sect.~\\ref{sec:results} the radiative transfer models and analysis of the observations are described.  Finally, we discuss the obtained results in Sect.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have observed several molecular species in comet \\machholz{} when it was close to perihelion in the sub-millimeter wavelength range with the SMT during six consecutive nights in January 2005.  These observations led to the detection of several \\detected{} rotational lines.  Our main goal was to measure the relative production rates of several parent and daughter volatiles in this comet and estimate the \\ce{CH3OH} rotational temperature.  We derive an HCN production rate of $(2.26 \\pm 0.03) \\times10^{26}\\ \\s$ at heliocentric distance of 1.2 AU, which corresponds to a $Q_\\ce{HCN}/Q_\\ce{H2O} \\sim 0.08$\\%, using a spherically symmetric radiative transfer numerical code that includes collisional effects between neutrals and electrons  and radiative pumping of the fundamental vibrational levels by solar radiation \\citep{1989A&A...216..278B}.  In addition to the statistical noise, the precision of the production rates is affected by the limited pointing accuracy at the time the observations were performed. An 8\\arcsec\\ pointing offset was included in the computation of the production rates, which results in $\\sim$  20\\% uncertainty in the derived values.  Mixing ratios relative to hydrogen cyanide and water are listed in Table~\\ref{tbl:mixing} for the production rates derived from the weighted average value of the observations on different dates during our campaign.  The asymmetric shape of the CO and CS line profiles suggests that there is preferential outgassing from the anti-sunward side of the nucleus, while the HCN and \\ce{H2CO} line profiles are fairly symmetric, although the latter is shifted toward the blue wing perhaps due to an instrumental effect. Using the rotational diagram technique, we retrieve a cold rotational temperature of \\rottemp\\ for the methanol energy levels sampled by the $J$ = 7--6 transitions.  Retrieved molecular abundances relative to water are comparable to those obtained in Oort Cloud comets \\citep{2002EM&P...90..323B,2009P&SS...57.1162C,2011ApJ...734L...8D}. Our $Q_\\ce{HCN}/Q_\\ce{H2O}$ mixing ratio of $\\sim 0.084$\\% is slightly lower than the typical value of 0.1\\% observed at radio wavelengths, and 50\\% lower than those found by \\citet{2009ApJ...699.1563B} from averaged 28 November 2004 and 19 January 2005 observations and by \\citep{2009ApJ...703..121K} from observations on 30 January 2005 using the same instrument.  The observed difference between our measurement of the HCN mixing ratio and those derived at infrared wavelengths is fairly typical.  The mixing ratio of \\ce{CH3OH} relative to \\ce{H2O} agrees with the infrared measurement in comet \\machholz{} by \\citet{2009ApJ...699.1563B} performed on 19 January 2005 at $\\rh = 1.208$ AU that is closer in time with our observations. Considering the pointing uncertainty that introduces an error of $\\sim$ 20\\%, our measurement is consistent within 1-$\\sigma$ with the revised abundances retrieved from the infrared observations on 19 and 30 January 2005 by \\citet{2012ApJ...747...37V} using a new line-by-line model for the $\\nu_3$ fundamental band of \\ce{CH3OH}.  The averaged CO production rate measured on 13 and 16 January agrees with that obtained by \\citet{2009ApJ...699.1563B} on 29 November 2004 at $\\rh = 1.493$ AU ($(6.3 \\pm 0.3)\\times10^{27} \\s$) within confidence limits.  In contrast, the $Q_\\ce{CO}/Q_\\ce{H2O}$ mixing ratio of $\\sim 2.6$\\% is almost a factor of two lower than that derived by \\citet{2009ApJ...699.1563B}.  For a parent molecule distribution, the derived $Q_\\ce{H2CO}/Q_\\ce{H2O}$ ratio is 0.14\\%, and $Q_\\ce{H2CO}/Q_\\ce{HCN}$ is 1.7\\%, which is more reliable since these molecules were observed simultaneously.  These values are intermediate between those measured by \\citet{2009ApJ...699.1563B} and \\citet{2009ApJ...703..121K}. On the other hand, the values inferred from a daughter molecule extended source distribution with a scale length of $L_\\mathrm{p} = 8000$ km are about a factor of 3 higher.  This illustrates that the derived production rates depend strongly on the assumed scale length of the parent molecule. This is particularly the case for \\ce{H2CO} where the beam size is slightly smaller than $L_\\mathrm{p}$, so it may be possible that this value overestimates the production rate.  An HNC abundance of 3.1\\% with respect to HCN is found assuming direct release from the nucleus with a Haser distribution (3.0\\% if infrared pumping of the fundamental vibrational levels is not considered), which suggests that this molecule may be destroyed by chemical reactions.  This value is compatible with those observed in other comets given the dependence on heliocentric distance -- except 73P/Schwassmann-Wachmann and the very active Hale-Bopp (C/1995 O1) \\citep[see][]{2008ApJ...675..931L}.  Outgassing variations induced by the nucleus rotation are expected to appear from non-sphericity of the nucleus or the presence of active region areas on the surface.  The variability observed in the HCN and CS outgassing rates on 12--13 January is affected by inaccurate pointing, which introduces an uncertainty of about the expected variation caused by the rotation of the nucleus.  HCN production rates show a uniform brightness increase over a period of two hours, which could be a periodic phenomenon consistent with the rotation period of the nucleus derived using different methods \\citep{2005IAUC.8480....3S,2007AJ....133.2001F,2009A&A...494..379R,2012Ap&SS.337..531M}, but our observations do not provide sufficient phase coverage to constrain the rotation period.  Comets are the most pristine objects in the solar system and have not undergone substantial thermal processing.  Cometary ices are more sensitive to thermal processing than dust. Hence their abundances provide indications about the formation and evolution of material in the early solar nebula. Dynamically new comets are expected to be enriched in volatile species, and CO is the most volatile component observed in \\machholz{}. However, our observations show a relative depletion of CO and and an intermediate-range mixing ratio of \\ce{H2CO} and \\ce{CH3OH} compared to the range of values measured in other comets \\citep{2002EM&P...90..323B}.  It is believed that the composition of icy material on the surface of the nucleus is altered by the exposure to solar radiation.  Thus, these observations provide a hint about its thermal/dynamical history and suggest that the comet has visited the inner solar system previously.  Formation regions of comets of Oort Cloud comets vary from 5--30 AU from the protosun according to the standard ``Nice model'' scenario \\citep{2008ssbn.book..275M} and can be constrained by observations of the chemical composition.  The relative depletion of several volatiles in \\machholz{} agrees with the determination of the formation region in the inner region of the solar nebula compared with other Oort Cloud comets, derived from the lower limit of the retrieved nuclear spin temperatures and dynamical models of the evolution of planetesimals in the solar system \\citep{2009ApJ...703..121K,2009ApJ...693..388K}.  Submillimeter spectroscopy of cometary atmospheres is a useful approach for studying the diversity of species that sublimate when a comet approaches the Sun.  These observations demonstrate the capabilities of the new CTS based on digital technology installed at the SMT to support solar system observation programs, further established with the cometary observations presented in \\citet{2004DPS....36.2505K,2004ExA....18...77V,2010A&A...510A..55D,2010ApJ...715.1258P,2011Icar..215..153J}, and observations of the Venusian mesosphere in \\citet{2008P&SS...56.1688R,2008P&SS...56.1368R}.  Our observations of comet \\machholz{} were analyzed for noise properties, standing wave and spectral line properties as compared to model predictions.  The spectrometer is found to perform according to its original specification in terms of sensitivity, spectral resolution and stability."
    },
    "1208/1208.4444_arXiv.txt": {
        "abstract": "Minor planet 2002 VE$_{68}$ was identified as a quasi-satellite of Venus shortly after its discovery. At that time its data-arc span was only 24 days, now it is 2,947 days. Here we revisit the topic of the dynamical status of this remarkable object as well as look into its dynamical past and explore its future orbital evolution which is driven by close encounters with both the Earth-Moon system and Mercury. In our calculations we use a Hermite integration scheme, the most updated ephemerides and include the perturbations by the eight major planets, the Moon and the three largest asteroids. We confirm that 2002 VE$_{68}$ currently is a quasi-satellite of Venus and it has remained as such for at least 7,000 yr after a close fly-by with the Earth. Prior to that encounter the object may have already been co-orbital with Venus or moving in a classical, non-resonant Near-Earth Object (NEO) orbit. The object drifted into the quasi-satellite phase from an L$_4$ Trojan state. We also confirm that, at aphelion, dangerously close encounters with the Earth (under 0.002 AU, well inside the Hill sphere) are possible. We find that 2002 VE$_{68}$ will remain as a quasi-satellite of Venus for about 500 yr more and its dynamical evolution is controlled not only by the Earth, with a non-negligible contribution from the Moon, but by Mercury as well. 2002 VE$_{68}$ exhibits resonant (or near resonant) behaviour with Mercury, Venus and the Earth. Our calculations indicate that an actual collision with the Earth during the next 10,000 yr is highly unlikely but encounters as close as 0.04 AU occur with a periodicity of 8 years. ",
        "introduction": "Minor planet 2002 VE$_{68}$ was discovered by Brian A. Skiff working for the LONEOS Survey at Lowell Observatory on November 11, 2002 and confirmed by the Eschenberg Observatory the following night (Griesser, Skiff \\& Spahr 2002)\\footnote{http://www.minorplanetcenter.org/mpec/k02/k02V52.html}. With a value of the semi-major axis $a$ = 0.7237 AU very close to that of Venus (0.7233 AU), this Aten asteroid is a Near-Earth Object (NEO) moving in a quite eccentric orbit, $e$ = 0.4104, that makes it a Mercury grazer, Venus crosser and Earth crosser. It has been designated a Potentially Hazardous Asteroid (PHA) by the Minor Planets Center (MPC) and as such has been the target of Doppler studies at Goldstone (Ostro \\& Giorgini 2004; Benner et al. 2008; Gavrik \\& Gavrik 2008) on November 2002 and 2010. These radar observations suggest that its near-surface is extremely rough. \\hfil\\par A preliminary rotational period of 13.5 h with a light curve amplitude $>$ 0.8 mag (indicating a very elongated body) and an estimated size of 260 m were found by Pravec, Wolf \\& Sarounova (2010)\\footnote{http://www.asu.cas.cz/$~$ppravec/neo.htm}. Bessel $BVRI$ photometry (Barajas et al. 2011)\\footnote{http://www.astronomerstelegram.org/?read=3073} has showed that 2002 VE$_{68}$'s mean colors are compatible with those of an X-type asteroid, perhaps similar to the E-type asteroid 2867 Steins (but also 1114 Lorraine, 5294 Onnetoh, 796 Sarita, 107 Camilla or 3686 Antoku). Barajas et al. (2011) also calculated a synodic period of 13.5 h (confirming the previous preliminary value), an albedo of about 0.25 and an absolute visual magnitude of 20.59 that gives an effective diameter of about 200 m (also consistent to preliminary determinations). With an amplitude of 0.9 mag, its light curve suggests that it may be a contact binary in which two rubble piles orbit a centre of mass in contact with each other (the full details of this research are available from the CURE at LACC web site\\footnote{http://www.lacitycollege.edu/academic/departments/physics/cure/reports/ BarajasT\\_Sp2011\\_Report.pdf}). This physical characterization is consistent with the battered surface suggested by radar data. \\hfil\\par Numerical computations by Mikkola et al. (2004) soon revealed that 2002 VE$_{68}$ is moving in a 1:1 mean motion resonance with Venus; more specifically, the asteroid is a quasi-satellite of Venus. As such, 2002 VE$_{68}$ is not a real, gravitationally bound satellite but from Venus point of view, the object appears to travel around it over the course of a Venusian year although it actually orbits the Sun. Venus has no known satellites: Sheppard \\& Trujillo (2009) completed a survey in search for satellites but no actual moons down to about 0.3 km in radius were detected. At the time of its identification as quasi-satellite of Venus, the arc length of 2002 VE$_{68}$ was only 24 days so that its orbit was not yet well known. The orbit has been improved significantly over the years and now it has an arc length of 2,947 days; besides, 2002 VE$_{68}$ has also been observed by radar (Ostro \\& Giorgini 2004; Benner et al. 2008). Here we revisit the topic of the current dynamical status of this remarkable object as well as look into its dynamical past and explore its future orbital evolution which is driven by close encounters with both the Earth and Mercury. In our calculations we use the most updated ephemerides and include the perturbations by the eight major planets. In addition, we include perturbations from the Moon and the three largest asteroids, (1) Ceres, (2) Pallas and (4) Vesta. \\hfil\\par This paper is organized as follows: in Section 2 the details of our numerical integrations are given. In Section 3, we present the results of our simulations. We discuss our results in Section 4. In Section 5 we compare our results with those obtained by Mikkola et al. (2004) and our conclusions are summarized in Section 6. ",
        "conclusions": "2002 VE$_{68}$, a remarkable NEO, was discovered by B. Skiff in 2002 (Griesser et al. 2002) and subsequently identified as a quasi-satellite of Venus (Mikkola et al. 2004). This paper revisits the dynamical status of 2002 VE$_{68}$, numerically integrating its trajectory using updated ephemerides and analyzing the results. We studied the orbit of 2002 VE$_{68}$ using different models and found good agreement between them on short time-scales. We can summarize the results of our investigation as follows:  a) We confirm the Venus quasi-satellite nature of 2002 VE$_{68}$ announced by Mikkola et al. (2004).  b) We confirm that 2002 VE$_{68}$ will leave its unusual dynamic status in a relatively short time-scale (about 500 yr).  c) We confirm that 2002 VE$_{68}$ got into its actual state after a close encounter with the Earth about 7,000-14,000 yr ago.  d) Close approaches are possible both at perihelion (with Mercury) and aphelion (with the Earth). Earth's are more important.  e) The influence of the Moon on the dynamics of 2002 VE$_{68}$ is not negligible as very close encounters with the Earth are possible.  f) 2002 VE$_{68}$ exhibits resonant (or near resonant) behaviour with Mercury, Venus and the Earth.  g) There is no danger of impact with the Earth, Venus or Mercury in the near future. Relatively close encounters with our planet have a periodicity of 8 years. The next close approach to the Earth will take place on November 4, 2018 at 0.038 AU.  Currently, the Earth has five known objects regarded as quasi-satellites: 3753 Cruithne (Wiegert, Innanen \\& Mikkola 1997), 2003 YN$_{107}$ (Connors et al. 2004; Brasser et al. 2004), (164207) 2004 GU$_{9}$ (Connors et al. 2004; Brasser et al. 2004; Wiegert et al. 2005; Mikkola et al. 2006; Wajer 2010), 2006 FV$_{35}$ (Mikkola et al. 2006; Stacey \\& Connors 2009; Wajer 2010) and 2010 SO$_{16}$ (Christou \\& Asher 2011). 2002 AA$_{29}$ will become a quasi-satellite of the Earth in the future (Connors et al. 2002). Venus has one, 2002VE$_{68}$ (Mikkola et al. 2004). Jupiter has four known quasi-satellites (Kinoshita \\& Nakai 2007): 2001 QQ$_{199}$, 2004 AE$_9$, P/2002 AR$_2$ LINEAR and P/2003 WC$_7$ LINEAR-CATALINA but new co-orbitals have recently been identified (Wajer \\& Kr\\'olikowska 2012). A few hundred Main Belt asteroids appear to be co-orbital (quasi-satellites in some cases) with (1) Ceres and (4) Vesta (Christou 2000; Christou \\& Wiegert 2012). Clearly the more objects identified in the quasi-satellite phase the better our understanding on their stability will be. Recognizing a variety of objects in the quasi-satellite state under different dynamical environments can only improve our knowledge on the overall processes that lead to the transformation of passing orbits into co-orbital ones and viceversa."
    },
    "1208/1208.4502_arXiv.txt": {
        "abstract": "\\noindent We report on \\chan\\, observations of the bright ultra--luminous X--ray (ULX) source in NGC~3921. Previous \\xmm\\, observations reported in the literature show the presence of a bright ULX at a 0.5--10 keV luminosity of 2$\\times 10^{40}$\\lum. Our \\chan\\, observation finds the source at a lower luminosity of $\\approx 8\\times10^{39}$\\lum, furthermore, we provide a \\chan\\, position of the ULX accurate to 0.7\\arcsec\\, at 90\\% confidence. The X--ray variability makes it unlikely that the high luminosity is caused by several separate X--ray sources. In 3 epochs of archival Hubble Space Telescope (HST) observations we find  a candidate counterpart to the ULX. There is direct evidence for  variability between the two epochs of WFPC2 F814W observations with  the observation obtained in 2000 showing a brighter source. Furthermore, converting the 1994 F336W and 2000 F300W WFPC2 and the 2010 F336W WFC3 observations to the Johnson $U$--band filter assuming a spectral type of O7~I we find evidence for a  brightening of the $U$--band light in 2000. Using the higher  resolution WFC3 observations the candidate counterpart is resolved  into two sources of similar color. We discuss the  nature of the ULX and the probable association with the optical  counterpart(s). Finally, we investigate a potential new explanation  for some (bright) ULXs as the decaying stages of flares caused by  the tidal disruption of a star by a recoiled supermassive black  hole. However, we find that there should be at most only 1 of such systems within z=0.08. ",
        "introduction": "Ultraluminous X-ray sources (ULXs) are off-nuclear X-ray point sources in nearby galaxies with X-ray luminosities, L${_X}\\approxgt 1\\times 10^{39}$ -- $10^{42}$ erg s$^{-1}$ (\\citealt{1999ApJ...519...89C}; \\citealt{2009Natur.460...73F}). Their X-ray luminosities are suggestive of intermediate-mass (10$^2$--10$^5$ \\msun) black holes (IMBHs) if they radiate isotropically at sub-Eddington levels as observed in stellar black holes and Active Galactic Nucleii (AGNs). Hence, ULXs could be a new kind of black hole with masses in--between the stellar--mass black holes found in X--ray binaries and the supermassive black holes (SMBHs; $\\approxgt 1\\times 10^6$ \\msun) found in the centers of galaxies. These IMBH could be the building blocks of SMBHs (e.g.~\\citealt{2010Natur.466.1049V}). For a comprehensive overview of ULXs see \\citet{2011NewAR..55..166F}.  Recent studies of the luminosity functions of ULXs and fainter extra-galactic X--ray binaries showed evidence for a break at a luminosity $\\sim2\\times 10^{40}$\\lum\\,(\\citealt{2011ApJ...741...49S}; \\citealt{2012MNRAS.419.2095M}). This, together with the identification of an ultra--luminous spectral state under super-Eddington accretion rarely observed in Galactic black hole X-ray binaries (e.g., \\citealt{2009MNRAS.397.1836G}), led to the idea that the majority of ULXs are stellar--mass black holes. These stellar--mass black holes have particular radiation mechanisms, including beaming effects (\\citealt{2001ApJ...552L.109K}) and/or truly super--Eddington emission (\\citealt{2002ApJ...568L..97B}). Nevertheless, a number of the brightest ULXs with an X--ray luminosity $\\approxgt 2\\times 10^{40}$\\lum\\, remains hard to explain as X--ray binaries. The higher the luminosity, the less likely it is that they can be explained as the high--luminosity end of the X--ray binary luminosity function.  A seemingly convincing case for an IMBH is provided by the variable, very bright ULX ESO~243--49 X--1 (\\citealt{2009Natur.460...73F}). Its X--ray luminosity is too high for a stellar--mass black hole even in the presence of some beaming. Given the evidence for the detection of a redshifted H$\\alpha$ emission line in the optical counterpart to ESO~243--49 X--1 (\\citealt{2010ApJ...721L.102W}), the uncertainty on its distance is reduced with respect to other bright ULXs. Another case for an IMBH is M82~X41.4+60, whereas it does not reach peak luminosities as high as ESO~243--49 X--1, the maximum luminosity is still uncomfortably high for a stellar--mass black hole (\\citealt{2003ApJ...586L..61S}).  \\citet{2009ApJ...696.1712F} showed that M82~X41.4+60 displays state changes similar to those of stellar--mass black holes when at sub--Eddington accretion rates. However, differences in interpretation exist. The temperature of the thermal emission is $> 1$ keV, which is high for an IMBH. \\citet{2006ApJ...652L.105O} describe the spectrum as that from a slim disk obtaining a black hole mass of 20--30\\,\\msun.  Similarly, \\citet{2009PASJ...61S.263M} interprete the thermal emission as coming from a thick corona and they find a black hole mass of $<$200 \\msun.  A further interesting alternative explanation for some (bright) ULXs is that the ULX is caused by a recoiling SMBH in the few million years after the recoil event (\\citealt{2010MNRAS.407..645J} and references therein). This scenario is not feasible for galaxies where there is evidence for the presence of a central supermassive black hole from e.g.~an AGN. The reason is that the recoiled SMBH needs to be replaced by a new SMBH in a subsequent merger, and the time for that is longer than the few million years lifetime of the recoiled SMBH as a ULX.  The potential explanation of the ULXs with luminosities above the break in the luminosity function as IMBHs or recoiled SMBHs warrant further investigation of these sources. So far, less than a dozen sources with ${\\rm L_X \\gtrsim 10^{41}}$\\lum\\, have been found (i.e.~in the Cartwheel galaxy, \\citealt{2006MNRAS.373.1627W}; M~82~X--1, \\citealt{2009ApJ...696.1712F} and \\citealt{2003ApJ...586L..61S}; ESO~243-49, \\citealt{2009Natur.460...73F}; in NGC~5775, \\citealt{2008MNRAS.390...59L}; and in CXO~J122518.6+144545, \\citealt{2010MNRAS.407..645J}). Eight more bright off--nuclear X--ray sources, of which 2 have ${\\rm L_X \\gtrsim 10^{41}}$\\lum, were recently reported by \\citet{2012MNRAS.423.1154S}, who suggest that their properties are consistent with accreting IMBHs in the hard state (cf.~\\citealt{2011MNRAS.416.1844W}).  Here, we present \\chan, archival Very Large Array (VLA), and archival HST/WFPC2 and HST/WFC3 observations of the bright ULX in \\src\\, first found by \\citet{2004MNRAS.353..221N} called NGC~3921~X--2. \\src\\, is classified as a protoelliptical that was formed during a major merger taking place approximately 0.7 Gyr ago, where one of the merging galaxies was gas rich and the other gas poor (\\citealt{1996AJ....111..109S}; \\citealt{1996AJ....112.1839S}). Given that \\src\\, has an AGN the recoiling black hole scenario desribed in \\citet{2010MNRAS.407..645J} is not applicable for the ULX under study in this galaxy. We adopt $\\Omega_m= 0.3$, $\\Omega_\\Lambda=0.7$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ to convert the redshift of \\src\\,to a distance measurement. ",
        "conclusions": "We report on a brief \\chan\\, observation of the field of the bright ULX in \\src\\ (\\citealt{2004MNRAS.353..221N}). Compared with the discovery \\xmm\\, observation $\\sim$9.5 years earlier, the source 0.5-10 keV X-ray luminosity has faded by a factor of $\\sim$2.5. This virtually proves that the high 0.5--10 keV X-ray luminosity of 2$\\times10^{40}$\\lum\\, detected by \\xmm\\, is due to a single bright source and not due to the chance super position of a group of sources. If the latter were the case they, coincidentally, would have decreased their X-ray luminosity by the time of the \\chan\\, observation (cf.~\\citealt{2007Natur.445..183M}).  Using the accurate position we obtained using \\chan\\, we investigate archival HST WFPC2 images in the F300W, F336W, F555W and F814W (2$\\times$) bands (which we converted to the Johnson $U$--, $U$--, $V$--, and Cousins $I$--bands, respectively) and archival WFC3 images in the F336W and F438W bands (which we converted to the Johnson $U$-- and $B$--band magnitudes).  There is a single bright, blue optical source in the astrometric error circle in the WFPC2 images. In the WFC3 images this source is resolved into two sources that, when combined, have the same F336W magnitude in the WFC3 and WFPC2 observations of 1994 and 2010 (but see below). With a F336W$-$F438W=$-0.14$, $U-B=0.3$, and $U-V=-0.2$ color, this source is blue (the color of source 1 and source 2 is approximately equally blue).  Furthermore, if at the distance of \\src, with a distance modulus DM of 34.5 excluding extinction effects, its combined absolute $V$-band magnitude, ${\\rm M_V}$, is $-10.6$. The Galactic extinction in the direction of \\src\\, is found to be $1.1\\times 10^{20}$ cm$^{-2}$ (\\citealt{1990ARA&A..28..215D}).  Using the relation of \\citet{1995A&A...293..889P} between ${\\rm N_H}$ and ${\\rm A_V}$ we obtain ${\\rm A_V}=0.06$. Using {\\sc synphot} we calculate the F336W-F438W color the O7~I star BZ~64 would have correcting for extinction. The observed F336W-F438W color is too blue to be explained by the O7~I star even if there was no extinction. The color of an O5~V star (BZ~77) may be consistent with the observed F336W$-$F438W color if the extinction is slightly lower than the E(B--V) of 0.18 derived from \\citet{1990ARA&A..28..215D}. So, we conclude that the source(s) making up source 1 and 2 have a combined spectrum similar to that of an early type star of approximately type O5.  A direct comparison of the two WFPC2 observations employing the F814W filter obtained with just over 6 years in between provides evidence for variability: in 1994 the F814W magnitude was 24.01$\\pm$0.11, whereas the observation in 2000 found F814W$=23.5\\pm 0.1$. The \\xmm\\, observation of \\citet{2004MNRAS.353..221N} was obtained on April 27, 2002 when the source was at its highest X--ray luminosity detected so far. The optical observation closest in time to the \\xmm\\, observation also finds the source at a slightly enhanced luminosity, suggesting that one of the two sources is indeed associated with the X--ray source.  Furthermore, comparing the Johnson $U$--band magnitudes calculated assuming an O7~I spectral type, the source is clearly brighter in 2000 than in 1994 and 2010, with $U$--band magnitudes of 23.7, 22.9 and 23.7 for the observation in 1994, 2000, and 2010, respectively, re-enforcing the evidence for variability.  The optical counterpart is classified as a stellar association or large globular cluster by \\citet{1996AJ....112.1839S} using the 1994 HST WFPC2 observations as we mentioned in Section~2.2 although they did not consider the $U$-band observations.  In the WFC3 observations the full--width at half maximum (FWHM) of source 1 and 2 is $\\approx 2.4$ pixels. This yields a FWHM of $<$0.1\\arcsec, which at the distance of 80 Mpc of \\src\\, is about 40 pc for each source. This is on the small side of the size distribution of OB associations (cf.~\\citealt{1998AJ....116..119B}) although the associations in IC~1613 have sizes of 40 pc (\\citealt{2010A&A...523A..23G}). The $U-V$ color is also consistent with that of young (10-50 Myr old) OB associations (\\citealt{2002A&A...392....1S}; \\citealt{2003A&A...401.1063A}). In a few nearby ULXs, e.g.~ NGC~1313~X--2 and Holmberg~IX~X--1, the age of the surrounding stars is found to be in the same range (\\citealt{2008A&A...486..151G}; \\citealt{2011ApJ...734...23G}). Similarly, assuming that one of the optical sources is related to the ULX and the other is an OB cluster the separation between them is similar to that for M~82~X--1 and NGC~7479 found by \\citet{2011MNRAS.418L.124V}.  Optical counterparts to ULXs with luminosities below $1\\times 10^{40}$\\lum\\, are typically fainter than the absolute magnitude for \\src\\,X--2, $M_V=-10.6$. The brightest have $M_V=-9$ (\\citealt{2008MNRAS.387...73R}).  The ULXs with higher luminosities are associated with brighter optical counterparts.  For instance, the bright ULX in ESO 243-49 has $M_R=-10$ (\\citealt{2010MNRAS.405..870S}) and the bright ULX in CXO~122518.6+144545 has $M_{g^\\prime}=-10.1$ (\\citealt{2010MNRAS.407..645J}). In the following we discuss the possible nature of this bright ULX and the associated potential counterpart.  \\subsection{\\src~X--2: a ULX with a stellar--mass black hole?}  The two candidate optical counterparts to \\src~X--2 are equally blue in the WFC3 observation. Since the accretion disk color can be similarly blue (e.g.~\\citealt{1995xrb..book...58V}), perhaps the observed light of one of the sources is dominated by the accretion disk. If so, the other source is dominated by a cluster of stars of a combined color similar to that of the accretion disk.  The observed X--ray and optical variability and in particular the large variability in the $U$--band does suggest that one of the two optical sources identified in the WFC3 images is indeed related to the counterpart to the ULX. If one of the two sources is due to a single object, it could also be a Luminous Blue Variable (LBV; \\citealt{1994PASP..106.1025H}). If there is evidence that this source is an LBV it is prudent to assume (Occam's razor) that it will be the mass donor to the accreting black hole in the ULX. The reason is that both LBVs and ULXs are so rare in a galaxy that the probability to find a ULX and an LBV very close together while not physically related is the product of the two (low) probabilities of finding one such source in a certain area in a galaxy.  On the other hand, one may be inclined to assume that evidence for an LBV for one of the two optical sources implies that the other source must be the ULX, reasoning that the number of LBVs in a binary orbit with a black hole must be rarer still than single LBVs. However, most O stars are formed in a binary with another (massive) star (\\citealt{2012Sci...337..444S}), making the scenario of finding an LBV with a black hole plausible.  Finally, a scenario where both blue optical sources detected in the WFC3 images are young OB clusters is also valid as long as one of them hosts the ULX. When the ULX is X--ray bright it has to have an optical counterpart of magnitude similar to that of the OB association, explaining the variability. Therefore, the scenario where the ULX in \\src~X--2 is explained as a stellar--mass black hole accreting from a donor in a wide system such that the absolute magnitude of the donor plus disk rivals that of an OB association is entirely consistent with the data.  \\subsection{\\src~X--2: a type IIn supernova?}  As the combined $U$--band magnitude from the WFC3 source 1 and 2 measured in 2010 is consistent with being the same as the WFPC2 $U$--band magnitude measured in 1994 while there is evidence for moderate brightening in 2000, a scenario explaining the ULX as a type IIn supernova (cf.~\\citealt{2009Natur.458..865G}; \\citealt{2010MNRAS.407..645J}) is strongly constrained. The scenario that still is (theoretically) possible but that we deem unlikely is as follows: the progenitor star of the supernova was not detected, therefore it should reside in a cluster that is still present. The supernova went off some time between 1994 and 2000 and, remarkably, went undetected in the optical. This includes a non-detection in archival groundbased observations from the Jacobus Kaptein Telescope in 1996. The supernova has faded in the optical before the HST observation in 2000 nearly down to the pre-explosion level. The X--rays come from the supernova and they are now fading.   \\subsection{A new scenario for the nature of some ULXs. Tidal disruption of a star by a recoiled black hole?}  A galaxy like our Milky Way is predicted to have hundreds of recoiled IMBHs/SMBHs with associated star clusters (\\citealt{2012MNRAS.421.2737O}).  For the majority of systems which are expected to have had recoil velocities $<$1000 km s$^{-1}$, the rate of tidal disruption events remains constant for about 1 Gyr after the recoil event. Given the expected rate one could have about a 1000 events in 1 Gyr (\\citealt{2008ApJ...683L..21K}).  Tidal disruption of stars from such a bound (nuclear) cluster or from unbound stars of the galaxy will produce periods of enhanced accretion providing X-ray luminosities above $10^{40}$ \\lum\\, (\\citealt{2008ApJ...683L..21K}; \\citealt{2012MNRAS.422.1933S}). The short duration Eddington--limited flares may well be missed by current scanning X-ray All Sky Monitors, whereas the flares may be too soft to be detected by the {\\it Swift} BAT instrument (cf.~\\citealt{2011MNRAS.410..359L} for the predicted lightcurves and the disk temperatures).  Theoretically, one expects the timescale of the decay of mass accretion rate to go as $t^{-5/3}$ (\\citealt{1988Natur.333..523R}; \\citealt{1990Sci...247..817R}). Recent \\chan\\, observations of three potential tidal disruption events are consistent with this theoretical trend and an overall lifetime of approximately 5 years at L$_x >10^{40}$ \\lum\\,(\\citealt{2004ApJ...604..572H}). In order to determine the importance of tidal disruption events from recoiled, off-nuclear black holes for the population of ULXs we need to estimate the number of such black holes that recoiled in approximately the last Gyr. The rate of tidal disruption events drops off steeply for longer times as the loss-cone of the recoiled black hole is emptied (\\citealt{2008ApJ...683L..21K}).  \\citet{2000ApJ...536..153P} found that over the last 1 Gyr about 1\\% of galaxies with an absolute $B$--band magnitude $-21 \\leq {\\rm M}_B \\leq -18$ went through a merger (cf.~their formula 32). This was derived evaluating the number of interacting galaxies found in the Second Southern Sky Redshift Survey. As the authors explain, this fraction is a lower limit as only pairs of relatively bright galaxies have been included. Including fainter galaxies would increase the fraction as more interacting bright $+$ faint galaxy pairs would be found. More recently, \\citet{2007ApJ...666..212D} found a merger fraction of 2\\% using close pair fractions and asymmetries from the Millenium Galaxy catalogue.  Assuming that the black hole merger time is significantly less than 1 Gyr such as in gas--rich mergers (e.g.~\\citealt{2007Sci...316.1874M}) will provide a lower limit on the number of recently (within the last Gyr) recoiled black holes. In the case that this merger time is much longer (of order of several Gyrs) the number of recently recoiled black holes is larger as the galaxy merger fraction was high at larger redshift (e.g.~\\citealt{2000ApJ...536..153P}; \\citealt{2004ApJ...617L...9L}).  In order to derive the number of off--nuclear recoiled black holes that recoiled in approximately the last Gyr we need to know the fraction of merging black holes that receive a significant kick. This fraction is debated at present. The magnitude of the recoil kick depends on the mass ratio of the two black holes before their merger, on the orientation of their spin axes with respect to the binary orbital plane and on the amount of spin. The presence of massive gas discs around the binary black hole such as potentially in gas--rich mergers can align the spin axes of the two black holes prior to their merger, reducing the amplitude of the recoil kick imparted during the merger (\\citealt{2007ApJ...661L.147B}; \\citealt{2009MNRAS.tmp.1795D}). However, fragmentation of the accretion disc around the binary SMBH could reduce the rate of alignment as fragments could result in short randomly oriented accretion events (\\citealt{2007MNRAS.377L..25K}). Recent calculations seem to suggest that the recoil kicks might be limited to velocities lower than $\\approx 200$ km s$^{-1}$ due to a gravitational wave \"anti-kick\" during the ringdown phase (\\citealt{2010CQGra..27a2001L}). Similarly, under certain conditions relativistic precession will align or anti-align the black hole spin vectors reducing the median kick velocities (\\citealt{2010ApJ...715.1006K}).  Recoil velocities of a few hundred km s$^{-1}$ imply that many recoiling SMBHs will be retained by the host galaxy. Too small recoil velocities will mean that the recoiled black hole will sink back to the nucleus of the galaxy and X-ray emission due to decaying tidal disruption events will not be identified as ULXs. For now we assume that 50\\% of the recoiled black holes will have recoil velocities such that they are bound to the galaxy and that they appear as off-nuclear sources for 1 Gyr. We note however, that this number depends on the distribution of recoil kick velocities which is uncertain.  Putting all these estimates together we calculate the probability of finding a ULX caused by tidal disruption of a star by a recoiled SMBH per galaxy. There is 2\\% (the merger fraction) $\\times$ 50\\% (right recoil velocity) $\\times$ $1\\times10^{-6}$/yr (the tidal disruption rate per year) $\\times$ 5 (the life--time in years above a luminosity of 1$\\times 10^{40}$\\lum) leading to about 5 in $10^8$ galaxies that potentially has a bright off--nuclear X--ray source due to the fading of emission caused by a tidal disruption event.  The 1 Gyr look--back time corresponds to a redshift of z$\\approx 0.08$.  This provides a co-moving volume of 0.16 Gpc$^{-3}$. Given the galaxy density $\\phi_\\star \\approx 1\\times 10^{-2}$ Mpc$^{-3}$ (cf.~\\citealt{1976ApJ...203..297S}; \\citealt{1992ApJ...390..338L}; \\citealt{2005MNRAS.360...81D} for h=0.7) there are about 1.6 million galaxies in this volume. Therefore, there will be between 0.1--1 X-ray source with L$>10^{40}$\\lum\\, due to tidally disrupted stars from off-nuclear, recoiled, black holes within a redshift of 0.08. Note that a luminosity of L$>10^{40}$\\lum\\, at a redshift of z=0.08 corresponds to a flux of only $\\approx 7\\times 10^{-16}$\\flx\\, making it difficult to detect these sources in the existing shallow surveys. If the estimates above are reasonable there might be at most one of these among the (bright) ULXs such as the ULX \\src~X--2.  \\subsubsection{Application to \\src~X--2}  The decay in X--ray luminosity of a factor of 2.5 in \\src~X--2 that we find between the \\xmm\\, and \\chan\\, observations $\\sim3400$ days apart is (too) small compared to the predicted drop-off from the soft disk X--ray spectral component predicted by \\citet{2011MNRAS.410..359L}. Even if in some proposed tidal disruption events the X--ray spectrum is found to be significantly harder than expected, the luminosity is about a factor 200 lower 3165~days after the soft peak (\\citealt{2004ApJ...603L..17K}). Given the presence of nuclear activity in \\src\\, a new central SMBH must have replaced the recoiled black hole, for instance following a recent merger, or the recoiled black hole was formed by a three black hole interaction event. All these constraints of course make the scenario of a tidal disruption event in a recoiled SMBH for \\src\\,X--2 unlikely.  \\subsection{Alternative explanations for \\src~X--2?}   Alternative explanations for the bright ULX in \\src\\, such as a background AGN or a foreground star or (quiescent) low--mass X--ray binary (LMXB) are possible in principle. For \\src~X-2 to be an active LMXB the distance should be $\\approx$1 Mpc for a luminosity of 1$\\times 10^{36}$\\lum\\, typical for an LMXB. This would give an absolute $V$--band magnitude of around -1 which is in line with that of LMXBs (\\citealt{1995xrb..book...58V}), however, finding an LMXB in interstellar space is unlikely. \\src~X--2 can also be a quiescent LMXB at a distance of a few kpc. The absolute magnitude of the optical counterpart would then be pointing to an M--type companion, which would imply that the blue candidate optical counterparts have to be unrelated to the quiescent LMXB.  \\citet{1996AJ....112.1839S} estimate the number of background galaxies in the WFPC2 field of view as 1.2$\\times 10^{-3}$ per arcsec$^{-2}$. The presence of the background galaxy cluster Abell 1400 is taken into account in this estimate. The optical counterpart does coincide spatially with one of the tidal tails that clearly belongs to \\src. As the (young) stellar associations in \\src\\, trace these tails very well the probability of finding a background AGN among these stellar associations is low.  In the case that the optical source is unrelated to the X-ray source the observed variability in the F814W (and $U$--band) WFPC2 observations is hard to explain.  Nevertheless, assuming the variability is spurious or explained in another way, we can provide a limit on the absolute magnitude of the counterpart of M$_U>-7.3$, M$_B>-8.2$, M$_V>-7.7$ and M$_I>-8.8$. These limits on M$_U$, M$_B$, M$_V$ and M$_I$ have been derived using the upper limits on the apparent magnitude of $<27.2$, $<26.3$, $<26.8$ and $<25.7$, respectively and given the distance of 80 Mpc towards \\src.  Therefore, the source could be a ULX with a stellar counterpart outside an association.  The upper limit to an optical source in this scenario translates to a lower limit on the ratio of the X--ray flux to the optical flux of 12.  Whereas most \\chan--selected AGN have a ratio of X--ray to optical $R$--band flux of lower than 10 (\\citealt{2009ApJS..180..102L}), our limit does not rule out a background AGN scenario in the case that the optical source is unrelated to the X--ray source. The most secure way to differentiate between the discussed scenarios is to obtain an optical spectrum of the potential counterparts."
    },
    "1208/1208.5990_arXiv.txt": {
        "abstract": "We report the discovery of optical filamentary and diffuse emission from G182.4+4.3 using 1.5-m Russian-Turkish telescope. We present the optical CCD images obtained with H$\\alpha$ filter revealing the presence of mainly filamentary structure at the northwest, filamentary and diffuse structure at the centre, south and north regions of the remnant. The bright optical filaments located in northwest and south regions are well correlated with the prominent radio shell of the remnant strongly suggesting their association. From the flux-calibrated CCD imaging, the average [S\\,{\\sc ii}]/H$\\alpha$ ratio is found to be $\\sim$0.9 and $\\sim$1.1 for south and northwest regions, which clearly indicates that the emission originates from the shock heated gas. We also present the results of X-ray data obtained from {\\it XMM-Newton} that show diffuse emission with a very low luminosity of $\\sim$$7.3\\times10^{31}$ erg s$^{-1}$ at a distance of 3 kpc in 0.3$-$10 keV energy band. Furthermore, we find a surprisingly young age of $\\sim$4400 yr for this remnant with such a large radius of $\\sim$22 pc. ",
        "introduction": "There are 274 Galactic supernova remnants (SNRs) catalogued in \\citet {b12} by their radio emission. Most of them are located in the Galactic plane where the density of gas and dust is very high which make the observation, except for radio band, of SNRs hard due to strong interstellar extinction and reddening effects in their line of sight. This difficulty can be surmounted by using narrow-band filters such as H$\\alpha$, [S\\,{\\sc ii}], [N\\,{\\sc ii}] and [O\\,{\\sc iii}] centered on characteristic emission lines on deep exposures. A number of Galactic SNRs are detected especially with H$\\alpha$ filter (e.g. \\citet {b9, b6, b8, b3, b4}), in which they usually have peculiar morphological structures. Some of them have filamentary structure while others have filamentary and diffuse structures together or arc (shell) structures. \\citet {b5} reported a catalogue of 24 known Galactic SNRs uncovered in H$\\alpha$ light in the Anglo-Australian Observatory/United Kingdom Schmidt Telescope (AAO/UKST) H$\\alpha$ survey of the southern Galactic plane. The optical observations of Galactic SNRs enable us to study the physical conditions in the remnant and the ambient medium, such as variation of chemical composition, density and evolutionary state.  Galactic SNR G182.4+4.3 located in the anti-centre region has been detected at 1400 MHz, 2675 MHz, 4850 MHz, and 10450 MHz with the Effelsberg 100-m telescope in radio bands by \\citet {b1}. They reported that the remnant has a shell structure with a radio spectral index of $\\alpha$=$-0.42\\pm0.10$ (S$_{\\nu}\\sim\\nu^{\\alpha}$), about 50 arcmin in size, very low radio surface brightness (7.5$\\times 10^{-23}$ Watt m$^{-2}$ Hz$^{-1}$ sr$^{-1}$ at 1 GHz), and is in Sedov expansion stage. Expanding into an ambient medium of low density indicated that its location was in front of the outer spiral arm III. H\\,{\\sc i} line observations of neutral hydrogen showed that the column density towards the remnant must be $\\leq$4$\\times$10$^{21}$ cm$^{-2}$. Considering its low ambient density and distance to be 3 kpc, they argued that z-height of the remnant was about 230 pc and concluded that it was most likely the remnant of a Type Ia supernova explosion. They have searched the {\\it ROSAT} all sky survey at the position of this remnant and detected no visible X-ray emission (corresponding upper limit of 6.2$\\times 10^{-2}$ counts s$^{-1}$). Assuming an initial explosion energy of $10^{51}$ erg, they obtained ambient density $n_{0}\\sim 0.013$ ${\\rm cm}^{-3}$, swept-up mass $M_{\\rm sw}\\sim 14$ M\\sun, electron temperature $kT\\sim 5.0$ keV, age 3800 yr, and shock velocity $V_{\\rm exp}=2300$ km ${\\rm s}^{-1}$.   There is no optical identification of G182.4+4.3 in literature so far. In this paper, we report the first detection of optical emission from G182.4+4.3 both filamentary and diffuse structure in H$\\alpha$. We also studied {\\it XMM-Newton} data to investigate its properties in X-ray bands and combine optical imaging and X-ray observation to present a better view of this SNR.  The paper is organized as follows: Observations and data reduction are described in Section 2. Based on the CCD imaging analysis together with X-ray results, we discuss the filamentary and diffuse structure and the plasma parameters of the SNR in Section 3. We give our conclusion in Section 4.   \\section[]{Observations and data reduction} \\subsection{Optical Imaging} Optical CCD imaging observations of G182.4+4.3 were taken on 2011 February 05, March 02 and 2012 February 19, with low resolution spectrograph TFOSC (TUG Faint Object Spectrograph and Camera) equipped with  2048$\\times$2048 back-illuminated camera with a pixel size of 15 $\\mu m$ $\\times$ 15 $\\mu m $ in a 13.5$\\times$13.5 arcmin$^{2}$ field of view (FOV) attached to the Cassegrain focus of the 1.5-m Russian-Turkish joint telescope (RTT150)\\footnote{Details on the telescope and the spectrograph can be found at http://www.tug.tubitak.gov.tr.} at T\\\"{U}B\\.{I}TAK National Observatory (TUG) T\\\"{u}rkiye, Antalya. The optical images were obtained with H$\\alpha$, [S\\,{\\sc ii}] and their continuum filters. The characteristics of these interference-filters are summarized in Table 1. The images were reduced by using standard {\\sc iraf} (Image Reduction Analysis Facility) routines for background subtraction, flat-fielding, trimming and continuum subtraction. The spectroscopic standard star HR5501 \\citep {b10, b11} was used for the absolute flux calibration. We write the coordinate information into the FITS header of the individual images with world coordinate system (WCS) tools. The weather conditions during these imaging observations were poor.  Due to its large size, we divided the whole remnant into several fragments and observed each fragment with H$\\alpha$ filter with short exposure time of 100 s. We obtained optical emission only from four regions, namely north (N), northwest (NW), center (C) and south (S). We focused our study on these four regions to obtain images with H$\\alpha$, [S\\,{\\sc ii}] and their continuum filters with relatively long exposures (900 s). Continuum subtracted H$\\alpha$ images of NW, S, C and N regions are given in Fig. 1, 2 and 3.  We also took the long-slit spectra by locating the slit on the brightest filament in S region on 2011 November 07 with an exposure time of 7200 s, with the spectroscopic mode of the TFOSC attached to Cassegrain focus of the RTT150. Unfortunately, due to poor observing conditions the obtained spectra did not allow us to get reliable results.  \\subsection{{\\bf X-ray and radio-continuum observations}} {\\it XMM-Newton} \\citep {b32} observed G182.4+4.3 on 2001 Jun 15, under the observation ID 50406540 and exposure time of 23 ks. {\\it XMM-Newton} has three X-ray telescopes \\citep {b33}, one equipped with EPIC-PN \\citep {b30} and two with EPIC-MOS \\citep {b31} CCD detectors in the focal plane. The data were reduced and analyzed using the {\\sc xmm-newton sas}\\footnote{Science Analysis Software ({\\sc sas}), http://xmm.vilspa.esa.es/sas/}, version 1.52.8. Calibrated event files for the EPIC-MOS1, EPIC-MOS2 and EPIC-PN detectors were produced using {\\sc sas} task {\\sc emchain} and {\\sc epchain} and following standard procedures. EPIC-MOS2 image in 0.3$-$10 keV energy band of the G182.4+4.3 is given in Fig. 4.  The overall radio structure of the G182.4+4.3 at 4850 MHz with the Effelsberg 100-m telescope (R. Kothes, private communication) is given in Fig. 5. As seen from the figure, the radio shell is most prominent in the south-west direction, this is because the prominent part of the shell is expanding towards the Galactic plane into higher density medium while the top shell is expanding away from the plane into very low density medium \\citep {b2}.  We overlaid the H$\\alpha$ mosaic image of the observed regions with radio-continuum contour image taken from 4850 MHz Effelsberg data and X-ray contour image from {\\it XMM-Newton} data to see if there is an association between the optical, radio and X-ray emission. As seen from the Fig. 6, there is a good correlation between optical and radio emission in NW and S regions. Due to poor observing conditions, for N and C regions the correlation is not seen clearly. The optical emission of S region correlates with both the radio and X-ray emission. N, NW and C regions are stand out of the FOV of {\\it XMM-Newton}. ",
        "conclusions": "We detected optical filamentary and diffuse emission for the first time for G182.4+4.3 with several short or long scale filaments found in the N, NW, C and S regions. The strong optical filaments located in NW and S regions match with the prominent radio shell at 4850 MHz Effelsberg data of the remnant suggesting their association. A very strong emission of [S\\,{\\sc ii}] relative to H$\\alpha$ [[S\\,{\\sc ii}]/H$\\alpha \\sim$0.9, $\\sim$1.1] obtained from imaging suggests that the emission originates from the shock heated gas. Finally, we estimated X-ray properties of this remnant using {\\it XMM-Newton} archival data. The X-ray spectrum of G182.4+4.3 shows that the plasma is of thermal origin in non-equilibrium ionization state and requires a temperature of $\\sim$0.9 keV. The best-fitting of the spectrum implies that the SNR is very young ($\\sim$4400 yr), expanding in a very low density medium ($\\sim$0.024 $\\rm cm^{-3}$), it has a low X-ray emitting mass ($\\sim$$1{\\rm M}\\sun$) and very low X-ray luminosity ($\\sim$$7.3\\times 10^{31}$ erg s$^{-1}$ in 0.3$-$10 keV)."
    },
    "1208/1208.0267_arXiv.txt": {
        "abstract": "{ It was recently reported that there may exist monochromatic $\\gamma$-ray emission at $\\sim 130$ GeV from the Galactic center in the Fermi Large Area Telescope data, which might be related with dark matter (DM) annihilation. In this work we carry out a comprehensive check of consistency of the results with the DM annihilation scenario, using the $3.7$ yrs Fermi observation of the inner Galaxy, Galactic halo, clusters of galaxies and dwarf galaxies. The results found are as follows. 1) Very strong constraints on the DM annihilation into continuous $\\gamma$-rays from the Galactic center are set, which are as stringent as the ``natural'' scale assuming thermal freeze-out of DM. Such limit sets strong constraint on the DM models to explain the line emission. 2) No line emission from the Galactic halo is found in the Fermi data, and the constraints on line emission is marginally consistent with the DM annihilation interpretation of the $\\sim 130$ GeV line emission from the inner Galaxy. 3) No line emission from galaxy clusters and dwarf galaxies is detected, although possible concentration of photons from clusters in $120-140$ GeV is revealed. The constraints from clusters and dwarf galaxies are weak and consistent with the DM annihilation scenario to explain the $\\sim 130$ GeV line emission.} ",
        "introduction": "Through analyzing the public Fermi Large Area Telescope (Fermi-LAT) data, several groups reported the hints of monochromatic $\\gamma$-ray emission with energy $E\\approx130$ GeV \\cite{2012JCAP...07..054B,2012JCAP...08..007W,2012JCAP...09..032T, 2012arXiv1205.4700B,2012arXiv1206.1616S}. The morphology of the potential line emission is still in debate. In some works the results showed that there were various regions lying basically in the Galactic plane having excess in $120-140$ GeV \\cite{2012JCAP...09..032T,2012arXiv1205.4700B}, while in another work the excess concentrated in the inner $5^{\\circ}$ of the Milky Way center \\cite{2012arXiv1206.1616S}. The results stimulated active discussions of the possible dark matter (DM) origin \\cite{2012arXiv1205.1520D, 2012PhRvD..86a5016C,2012PhRvD..86d3515C,2012arXiv1205.4151K, 2012arXiv1205.4675L,2012JCAP...09..003R,2012arXiv1205.5789S, 2012PhRvD..86h3521C,2012arXiv1206.2910W,2012arXiv1206.4758F, 2012arXiv1207.1341H,2012arXiv1207.1537O,2012PhLB..715..285Y, 2012JCAP...10..033F,2012arXiv1207.4981P,2012JCAP...08..003D, 2012arXiv1206.2863K}. Alternatively the astrophysical explanations of the line-like $\\gamma$-ray emission were also proposed \\cite{2012JCAP...07..011P,2012arXiv1207.0458A}.  Due to the importance of the high energy line emission, many works were trying to test the line emission (and its DM interpretation) with Fermi-LAT observations of continuous $\\gamma$-rays \\cite{2012PhRvD..86d3524B,2012arXiv1206.7056B,2012arXiv1207.0800C, 2012PhRvD..86h3525C}, the line emission from dwarf galaxies and clusters of galaxies \\cite{2012PhRvD..86b1302G,2012arXiv1207.4466H}, the unassociated Fermi sources \\cite{2012arXiv1207.7060S}, and the future detection with high energy resolution telescope \\cite{2012PhLB..715...35L,2012arXiv1207.6773B}. The basic results show that the allowed cross section of the DM annihilation final states giving rise to the continuous emission can only be about $O(10^2)$ times larger than that of the line emission \\cite{2012PhRvD..86d3524B,2012arXiv1206.7056B,2012arXiv1207.0800C, 2012PhRvD..86h3525C}, which needs to be considered when constructing the DM models. The recent search for line emission in the Milky Way halo by Fermi collaboration with two-year PASS 6 data showed no indication of signals \\cite{2012PhRvD..86b2002A}, and the upper limits seemed to be marginally in conflict with the results found in \\cite{2012JCAP...08..007W,2012JCAP...09..032T}. Furthermore, no signal from dwarf galaxies was found and the constraints on DM annihilation into monochromatic $\\gamma$-rays were consistent with the results found in the inner Galaxy \\cite{2012PhRvD..86b1302G}. For the galaxy clusters, however, it was claimed a $\\sim3\\sigma$ signal of possible $130$ GeV line emission by Hektor et al. \\cite{2012arXiv1207.4466H}.  In this work we try to make a comprehensive test of the DM scenario of the line emission, through analyzing the Fermi-LAT $\\gamma$-ray data in the Galactic center, the Milky Way halo, dwarf galaxies and galaxy clusters. In our analysis the information of the spatial distribution of DM is taken into account. We further include the effects of substructures on the DM annihilation in the analysis, both the enhancement of annihilation luminosity and the change of surface brightness distribution \\cite{2012MNRAS.419.1721G,2012MNRAS.425.2169G}.  This paper is organized as follows. In Sec. II we will discuss the constraints on the continuous $\\gamma$-ray emission from the Galactic center region. In Sec. III, IV and V we give the results of the line search from the Milky Way halo, galaxy clusters and dwarf galaxies respectively. Finally Sec. VI is our conclusion and discussion. ",
        "conclusions": "The recently reported tentative $\\gamma$-ray line emission from the Fermi-LAT observations in the inner Galaxy \\cite{2012JCAP...07..054B, 2012JCAP...08..007W,2012JCAP...09..032T,2012arXiv1205.4700B, 2012arXiv1206.1616S} has invoked many discussions of the DM annihilation signals. In this paper we do a comprehensive analysis using the Fermi-LAT data in the Galactic center region, Galactic halo, galaxy clusters and dwarf galaxies, to test the DM annihilation interpretation of this line emission.  Using the Fermi-LAT data in the Galactic center region, we get strong constraints on the continuous $\\gamma$-ray emission of DM annihilation to $W^+W^-$, $b\\bar{b}$, $\\mu^+\\mu^-$ and $\\tau^+\\tau^-$ final states, which set useful constraints on DM models to explain the $130$ GeV line emission.  We further perform the search for line emission with Fermi-LAT data in the Milky Way halo. The constraints on the line emission are generally consistent with the tentative ``signal'' found in the inner Galaxy \\cite{2012JCAP...08..007W}. Only when the enhancement effect due to DM substructures is taken into account, and the ``CLEAN'' events are studied, the constraints seem to have a weak tension with the DM annihilation interpretation of the $130$ GeV line emission if the DM density profile follows Einasto or NFW profiles. However, considering the uncertainties of both the constraints and the fitting ``signal'', we can not draw a firm conclusion on it at present. We just mention that if it is finally proven to be true, one may need to assume cuspier density profile of DM in the Galactic center (e.g., accreted by the central supermassive black hole \\cite{1999PhRvL..83.1719G}) to consistently understand the results in the inner Galaxy and the Galactic halo within DM annihilation scenario.  Galaxy clusters and dwarf galaxies are also studied to further test the DM interpretation of the $130$ GeV line emission. It is interesting to find that there is a concentration of photons in $120-140$ GeV from the six nearby clusters (Figure 4) for the ``CLEAN'' and ``ULTRACLEAN'' classes of events. However, the result seems hard to indicate an excess of line emission. Therefore we set upper limits on the DM annihilation cross section to monochromatic $\\gamma$-ray line, which are consistent with the claimed results in the inner Galaxy \\cite{2012JCAP...08..007W}, even when the large boost factor of substructures are taken into account. No signal of line emission from dwarf galaxies is found, and the constraints on DM annihilation cross section to monochromatic $\\gamma$-ray line are derived. Compared with the galaxy clusters, the constraints are weaker for dwarf galaxies, and are consistent with the claimed results in the inner Galaxy \\cite{2012JCAP...08..007W}."
    },
    "1208/1208.4787_arXiv.txt": {
        "abstract": "We consider the three-dimensional instability of a layer of horizontal magnetic field in a polytropic atmosphere where, contrary to previous studies, the field lines in the initial state are not unidirectional. We show that if the twist is initially concentrated inside the unstable layer, the modifications of the instability reported by several authors (see e.g. \\citet{cattaneo1990}) are only observed when the calculation is restricted to two dimensions. In three dimensions, the usual interchange instability occurs, in the direction fixed by the field lines at the interface between the layer and the field-free region. We therefore introduce a new configuration: the instability now develops in a weakly magnetised atmosphere where the direction of the field can vary with respect to the direction of the strong unstable field below, the twist being now concentrated at the upper interface. Both linear stability analysis and non-linear direct numerical simulations are used to study this configuration. We show that from the small-scale interchange instability, large-scale twisted coherent magnetic structures are spontaneously formed, with possible implications to the formation of active regions from a deep-seated solar magnetic field. ",
        "introduction": "Active regions at the surface of the Sun are believed to be the visible manifestation of deep-seated intense magnetic fields. The generation of these predominantly toroidal fields is associated with the differential rotation in the tachocline. However, the transport of these strong fields from the lower part of the convective zone up to the photosphere remains one of the major unknowns of the solar cycle. It is clear that super-equipartition fields are buoyant and could therefore rise up to the surface, but the observed size, coherence and twist of the magnetic structures forming active regions remain largely unexplained.  In order to obtain some understanding of the transport of magnetic structures by buoyancy, highly idealised models have been considered. An early model \\citep{parker1955} considers the buoyancy properties of the magnetic field when it is concentrated in slender pressure-confined flux tubes. A simple model describing the dynamical evolution of such structures was proposed by \\citet{spruit1981} and was subsequently used by several authors \\citep[see for example][]{moreno1983}. This model has successfully reproduced Joy's law  for the tilt of bipolar regions at the solar photosphere \\citep{choudhuri1987,dsilva1993,caligari1995} as well as the relationship between the tilt angle and the total flux of active regions \\citep{fan1994}. Many authors have sought to move on from this simple model and have considered the full set of compressible magnetohydrodynamic (MHD) equations (or the anelastic approximation) to study the evolution of a simplified magnetic structure mimicking a flux tube. These models have shown that a certain amount of twist must be present inside the tube in order for it to remain coherent as it rises; such a twist corresponds to a non-zero axial current inside the flux tube. Without twist, the flux tube is shredded apart by the vorticity generated by its own motion \\citep{schussler1979,longcope1996,emonet1998,fan1998,jouve2007}. Even more problematic than the issue of tube coherence is the interaction between the rising structures and the small-scale convective motions. Some studies have considered the rise of gradually twisted magnetic flux tubes through the convective zone both in Cartesian \\citep{fan2003,abbett2004} and spherical \\citep{jouve2009} geometries. Again, a given amount of twist is necessary for the flux tubes to remain coherent and rise through the convective zone. There is observational evidence of non-potential magnetic fields in active regions \\citep{seehafer1990,leka1996,pevtsov2001}, with the implication that the flux tubes are twisted before they emerge as active regions. While the twist is no more than one turn in general \\citep{chae2005}, the corresponding Lorentz forces are implicated in chromospheric flux eruptions \\citep{torok2005}. The origin of the twist is still unclear as there are various different explanations. The effect of the Coriolis force \\citep{fan2000} and  differential rotation \\citep{devore2000} on the rising structures can only contribute a small fraction of the observed twist values \\citep{holder2004}.  Other mechanisms have been suggested to explain the origin of non-potential magnetic fields. One possibility is that the tube becomes twisted due to the effect of small-scale helical motions in the convective zone acting on the rising magnetic field, the so-called $\\Sigma$-effect \\citep{longcope1998}. Another suggestion, which is closely related to the content of this paper, is that the twist arises by accretion of poloidal fields during the rise of a flux tube \\citep{choudhuri2003,chatterjee2006}. However indirect observational evidence suggests the presence of an initial twist in flux tubes at the base of the solar convective zone \\citep{holder2004}, so that the twist may be produced during the formation process of the flux tube itself.  Another issue, which we don't yet fully understand, is the mere existence of localised magnetic structures similar to flux tubes below the convective zone. What is known is that the differential rotation profile produces a strong toroidal field at the base of the solar convective zone. We can attempt to model the appearance of isolated flux structures by looking at a uniform layer of toroidal field where the vertical structure of the magnetic layer is initially imposed; then the horizontal length scale will be set naturally by the buoyancy instability. As a result of the magnetic pressure, the density of polytropic atmosphere in the magnetised region is reduced, assuming the temperature is not modified by the presence of the magnetic field. Consequently, the upper interface of the magnetic slab is susceptible to Rayleigh-Taylor-type instabilities. This paper is principally concerned with the three-dimensional evolution of these instabilities when the initial magnetic field has non-zero twist. The two-dimensional evolution of states with no initial twist was investigated by \\citet{cattaneo1988}; since no variation was permitted along the original field direction this treatment could only encompass the interchange instability, with the field lines remaining straight. This configuration has also been considered numerically in three dimensions by several authors \\citep{matthews1995,wissink2000}, whereas alternative initial conditions (in which the temperature profile is altered by the magnetic field instead of the density) were considered by \\citet{fan2001}. It has proved very difficult to produce flux structures of significant size from this type of instability. Whatever the scale of the original instability secondary vortex instabilities due to the down flows between the plumes disrupt the magnetic structures leaving a rather diffuse and dynamically inactive field . Note that an alternative model has been suggested by \\citet{kersale2007}, where they considered the non-linear evolution of magnetic buoyancy instabilities resulting from a smoothly stratified horizontal magnetic field. \\cite{cattaneo1990} conducted two-dimensional simulations of an initial twisted field structure. They found that in their constrained geometry large structures could appear as result of the twist. We will show in this paper that this effect disappears if we allow the motion to be fully three-dimensional.  The earlier simulations to examine the evolution of buoyant magnetic structures assumed that the region above the initial magnetic slab was field free. In this case the initial potential energy contained in the initial condition is eventually transferred to kinetic energy and is finally dissipated. Any possible effects of the magnetic field existing above the unstable slab are neglected. Is this a reasonable simplification?  As shown by many simulations of overshooting convection in the presence of magnetic field (see for example \\citet{tobias2001}), the strong toroidal field initially injected in the system or generated by a shear flow mimicking the tachocline \\citep{guerrero2011} is predominantly contained in the stable radiative zone. However, despite the effects of turbulent pumping, a non-negligible part of the field is still present in the convection zone, both due to the redistribution of large-scale magnetic fields and to local small-scale dynamo action. If the strong toroidal field is buoyantly unstable, it will rise and interact with the small scale magnetic perturbations present in the convective zone.  The purpose of this paper is to investigate the simplest model that allows us to consider the buoyancy instability of a toroidal magnetic field in a weakly magnetised atmosphere above. The layer of strong toroidal field models the field induced by shear in the tachocline whereas the magnetised atmosphere above represents the magnetic perturbations in the convective zone. For this first, illustrative study, the field in the convective zone is assumed to be unidirectional but with a different direction as in the layer of strong field below. This is of course highly idealised, but allows us to derive a model involving few free parameters. The instability of a sheared magnetic field has already been studied (\\citet{cattaneo1990,cattaneo1990b,kusano98,nozawa2005}). In this case, the twist is uniformly distributed inside the unstable slab and the atmosphere above is field-free. To our knowledge, this type of instability has only been considered in published papers using two-dimensional numerical simulations. The configuration we consider later in the paper is different as the field lines are unidirectional except in a thin region where the twist and current are concentrated. A similar setup was considered by \\citet{stone2007a,stone2007b} where they look at the effect of a transverse magnetic field on the Rayleigh-Taylor instability.  The paper will proceed as follows: in section \\ref{sec:model}, we describe the governing equations, the model and physical parameters and the numerical methods. We then extend the results by \\citet{cattaneo1990} and \\citet{kusano98} to the three-dimensional case in section \\ref{sec:shear}. Finally, the interchange instability in a weakly magnetised atmosphere is considered in section \\ref{sec:atmo}, using both linear stability analysis and non-linear numerical simulations in three dimensions. ",
        "conclusions": "}  In this work, we sought to examine the instability of a layer of horizontal magnetic field in a polytropic atmosphere, where the direction of the field lines is depth-dependent. In section \\ref{sec:shear} we examined the instability of a layer of horizontal magnetic field in a polytropic atmosphere, where the direction of the field lines in the layer varied with depth. We showed that the initial idea of building large-scale coherent magnetic structures from the buoyancy instability of a sheared magnetic layer, as studied by \\cite{cattaneo1990} and \\cite{kusano98}, seemed limited to two-dimensional geometry. The same configuration drastically changed in three dimensions, where the interchange instability was able to occur in an oblique direction. We stress that it could be possible to find a configuration of this type that strongly modifies the nature of the magnetic buoyancy instability, however this was not the aim of the paper.  In section \\ref{sec:atmo} we introduced another initial configuration, where the twist was concentrated at the interface between a strong layer of horizontal magnetic field and a weakly magnetised atmosphere above. The weakly magnetised atmosphere was justified by the presence of convection, which is expected to contain non-negligible magnetic fluctuations, both due to the redistribution of large-scale magnetic fields, and due to local small-scale dynamo action. The presence of a strong unidirectional field was justified by the presence of radial differential rotation in the tachocline. We presented a linear stability analysis and numerical simulations of the magnetic buoyancy instability in sections \\ref{sec:linear} and \\ref{sec:nonlinu3d}. We have shown that the initial interchange instability is only weakly modified by the presence of the transverse magnetic field. During the non-linear phase of the instability, the small-scale magnetic structures merged together to overcome magnetic tension, leading to the formation of magnetic structures on scales larger than those of the initial linear instability. These structures were shown to be more coherent than those in the unidirectional case, since the production of vorticity by the secondary instability was reduced. In addition we found a significant amount of twist was generated due to the rise of the toroidal field in the weakly magnetised atmosphere. This mechanism could provide the initial twist necessary for the magnetic structures to rise coherently through the convective zone.  It is however important to stress the limitations of the present model. In order to limit the number of parameters, we have considered two unidirectional fields with a current sheet at the interface. One could argue that the magnetic field in the convective zone is interacting with small-scale motions and is certainly not unidirectional as assumed here. Although it remains to be verified using a more refined model, we expect that locally the toroidal field will behave similarly in the presence of small-scale magnetic fluctuations.  There are two stabilising effects in the present model, the stably stratified atmosphere (which is necessary to suppress convective motions) and the magnetic tension generated by the twisted atmosphere. The tension becomes important as the atmosphere is compressed by the rising structures, and the toroidal magnetic field will eventually stop rising. We have reached this regime in some of our simulations where the pitch angle is equal or larger that $\\pi/4$. The stabilising effect of the tension is clearly overestimated by the present model, as the tension exists everywhere at the interface between the rising layer and the atmosphere. In a more realistic situation, the convective zone would be filled with small-scales magnetic structures. Locally these structures could have a similar effect to the global tension considered in our simulations, although we didn't check it numerically here. We plan in a forthcoming paper to look at a configuration where the global magnetic tension is removed at some depth, leading the rise and formation of local coherent twisted magnetic structures. Another possibility would be to extend the work of \\citet{fan2003} and \\citet{jouve2009}, assuming that the convective zone is filled with magnetic perturbations. This would confirm whether the generation of twist, via the interaction of a rising flux tube and small-scale magnetic perturbations, is sufficient for the flux tube to remain coherent as it rises through the convective zone. Such simulations are underway in spherical geometry using the ASH code \\citep{pinto2012}.  We conclude by arguing that the merging mechanism illustrated by this simple model could have implications on the formation of large-scale magnetic structures from a deep-seated toroidal field. The twist observed in active regions could therefore be attributed to a progressive accumulation of local twist as the toroidal field rises through the magnetised convective zone.  {\\bf Acknowledgements} This work has been financially supported by STFC. CPU time was provided by the HPC resources of CALMIP under the allocation 2012-P1118 and by the UKMHD supercomputing facility located in Warwick."
    },
    "1208/1208.6327_arXiv.txt": {
        "abstract": "Rastall's theory is a modification of General Relativity, based on the non-conservation of the stress-energy tensor. The latter is encoded in a parameter $\\gamma$ such that $\\gamma = 1$ restores the usual $\\nabla_\\nu T^{\\mu\\nu} = 0$ law. We test Rastall's theory in cosmology, on a flat Robertson-Walker metric, investigating a two-fluid model and using the type Ia supernovae Constitution dataset. One of the fluids is pressureless and obeys the usual conservation law, whereas the other is described by an equation of state $p_x = w_x\\rho_x$, with $w_x$ constant. The Bayesian analysis of the Constitution set does not strictly constrain the parameter $\\gamma$ and prefers values of $w_x$ close to $-1$. We then address the evolution of small perturbations and show that they are dramatically unstable if $w_x \\neq -1$ and $\\gamma \\neq 1$, i.e. General Relativity is the favored configuration. The only alternative is $w_x = -1$, for which the dynamics becomes independent from $\\gamma$. ",
        "introduction": "The nature of dark matter (DM) and dark energy (DE) is one of the most important open questions today in physics. There are strong observational evidences indicating that about $95\\%$ of the universe is under the form of DM ($\\approx 25\\%$) and DE ($\\approx 70\\%$), but no direct detection has been reported until now. The usual candidates to DM (e.g. neutralinos and axions) and DE (e.g. cosmological constant, quintessence) lead to very robust scenarios, but at same time they must face theoretical and observational issues. For recent reviews on the subject, see for example \\cite{Li:2011sd, Caldwell:2009ix, Bertone:2004pz}. \\par Among the alternatives to the standard description of the dark sector there is the possibility of a modification of gravity theory on large scales, see \\cite{Clifton:2011jh} for a recent review on the subject. An example are the so-called $f(R)$ theories, which are based on the inclusion of non-linear curvature terms in the Einstein-Hilbert action \\cite{DeFelice:2010aj}. Other possibilities are the Unified Dark Matter models, where DM and DE are treated as a single entity \\cite{Kamenshchik:2001cp, Gorini:2007ta, Piattella:2009da, Piattella:2009kt, Bertacca:2010mt, Campos:2012ez}, models in which DM is treated as a viscous component \\cite{Zimdahl:1996ka, Colistete:2007xi, HipolitoRicaldi:2010mf, Piattella:2011bs} and models of interaction in the dark sector (i.e. exchange of energy between DM and DE) \\cite{Zimdahl:2001ar, Zimdahl:2005bk}. \\par Another example, the one we pursue in this paper, is to touch one of the cornerstone of the gravity theory: the usual conservation laws for matter components. This kind of formulation was introduced by Rastall some 40 years ago \\cite{Rastall:1973nw, Rastall:1976uh}, and has been recently investigated in a cosmological context, giving some interesting results concerning the dynamics of the dark sector \\cite{Fabris:2011rm, Batista:2011nu, Fabris:2011wz, Daouda:2012ig, Fabris:2012hw}. \\par Rastall's motivation for modifying the usual conservation laws is based on the fact that the latter have been directly tested only locally or in a weak-field regime. On the other hand, the introduction of covariant derivatives imply, in some sense, an exchange of energy between matter and the gravitational field. Hence, in general, non-trivial generalizations of the conservation law are in principle possible. Besides, particle production in curved spacetimes is a central issue of quantum field theory on such spaces \\cite{Birrell:1982ix}. Therefore, one may regard modifications of the usual (classical) conservation laws as effective, semi-classical approaches to such phenomenon. \\par Rastall's proposal is the following: \\begin{equation} {T^{\\mu\\nu}}_{;\\mu} = \\kappa R^{;\\nu}\\;, \\end{equation} where the semicolon denotes the covariant derivative and $\\kappa$ is a (dimension-full) free parameter. The above relation can be rewritten as \\begin{equation}\\label{Rastmod} {T^{\\mu\\nu}}_{;\\mu} = \\frac{\\gamma - 1}{2} T^{;\\nu}\\;, \\end{equation} where $T$ is the trace of the stress-energy tensor and $\\gamma$ is now a dimensionless free parameter. When $\\kappa = 0$, then $\\gamma = 1$ and the usual conservation law (and thus General Relativity) is recovered. \\par In a one-fluid model, it is possible to redefine the energy-momentum tensor in order to recover the usual conservation law \\cite{Fabris:1998hr}. In this sense, Rastall's theory is just a redefinition of the fluid equation of state. However, in a multi-fluid case the modification introduced by Rastall opens possibilities of non-trivial interactions among the different components. In \\cite{Batista:2011nu}, this feature is used to investigate the consequences of the modified conservation law for a model of the dark sector of the universe: the authors investigate a two-fluid model, one of them being pressureless matter $p_{m} = 0$, whereas the other obeying the vacuum energy equation of state $p_{x} = - \\rho_{x}$. Assuming that the matter component obeys the usual conservation law, then the vacuum energy conservation law is affected by the presence of matter, via Eq.~\\eqref{Rastmod}. The main result of \\cite{Batista:2011nu} indicates that the model is completely equivalent to the $\\Lambda$CDM at the background and linear perturbations levels. There is just one striking difference: DE may now agglomerate. This fact could have important consequences at the non-linear level, which is a regime where the $\\Lambda$CDM faces some difficulties \\cite{Perivolaropoulos:2008ud}. \\par In the model studied in \\cite{Batista:2011nu}, the equivalence with the $\\Lambda$CDM at background and linear perturbations levels implies that no constraints on the parameter $\\gamma$ can be established using the corresponding observational tests. Such constraints are, on the other hand, in principle possible using non-linear data. In the present paper our goal is to verify to what extent the model studied in reference \\cite{Batista:2011nu} is a favorable configuration. To do this, we repeat the analysis made there but with the $x$ component now being described by a more general equation state, i.e. $p_{x} = w_x\\rho_{x}$, with $w_x$ constant. \\par We show that in this case the parameter $\\gamma$ appears explicitly in the background and linear perturbation equations. Therefore, using type Ia supernovae data, we calculate its probability distribution function (PDF) along with the one for the matter density parameter $\\Omega_{m0}$, and for the equation of state parameter $w_x$. The analysis show that supernovae date do not constrain strictly $\\gamma$ whereas $w_x \\sim -1$ is favored. Then, considering the evolution of small perturbations we find a dramatic instability if $w_x \\neq -1$ and $\\gamma \\neq 1$. This is a result which favors General Relativity, i.e. $\\gamma = 1$. On the other hand, another possibility seems to be viable, i.e. $w_x = -1$, for which the dynamical equations turn independent of $\\gamma$ and which is the case investigated in \\cite{Batista:2011nu}, which thus seems to be favored. \\par The paper is organized as follows. In Sec.~\\ref{Sec:RastThe} we present Rastall's theory, deriving some cosmological considerations. In Sec.~\\ref{Sec:Obsconstr} we confront the predictions of Rastall's theory with data from the type Ia supernovae of the Constitution set. In Sec.~\\ref{Sec:Pertan} we tackle the question of the evolution of small perturbations and, finally, in Sec.~\\ref{Sec:ResDisc} we present our conclusions. ",
        "conclusions": "\\label{Sec:ResDisc}  In this work we investigate in some detail Rastall's theory in cosmology, on a flat Robertson-Walker metric, investigating a two-fluid model where one of the fluids is pressureless and obeys the usual conservation law (therefore scales as $a^{-3}$), whereas the other is described by an equation of state $p_x = w_x\\rho_x$, with $w_x$ constant. This model is a generalization of the one studied in \\cite{Batista:2011nu}, for which $w_x = -1$ and for which the equivalence with the $\\Lambda$CDM at background and linear perturbations levels implies that no constraints on the parameter $\\gamma$ can be established using the corresponding observational tests. \\par We perform a Bayesian analysis on the type Ia supernovae Constitution dataset, and show that $\\gamma$ is not strictly constrained whereas $w_x \\sim -1$ is favored. On the other hand, considering the evolution of small perturbations we find a dramatic instability if $w_x \\neq -1$ and $\\gamma \\neq 1$. This is a result which favors General Relativity, i.e. $\\gamma = 1$. However, another possibility seems to be viable, i.e. $w_x = -1$, for which the dynamical equations turn independent of $\\gamma$. This is the case investigated in \\cite{Batista:2011nu}, which thus seems to be the only possibility in which a hydrodynamical model could work in Rastall's cosmology. We expect different predictions involving $\\gamma$ at the non-linear level of perturbations. A possibility to evade the strong constraints found in the perturbative analysis performed in the present paper is to describe the dark energy component through a self-interacting scalar field. The non-linear analysis and the scalar field formulation of the Rastall's cosmological model are projects for future researches."
    },
    "1208/1208.4608_arXiv.txt": {
        "abstract": "{The ubiquity of filaments in star forming regions on a range of scales is clear, yet their role in the star formation process remains in question.  We suggest that there are distinct classes of filaments which are responsible for their observed diversity in star-forming regions.  An example  of a massive molecular filament in the Galactic mid-plane formed at the intersection of UV-driven bubbles which displays a coherent velocity structure ($<$ 4 km s$^{-1}$) over 80 pc is presented.  We classify such sources as Massive Molecular Filaments (MMFs; M $\\geq 10^{4}$ \\Msun, length $\\geq$ 10 pc, $\\bigtriangleup$v $\\leq$ 5 km s$^{-1}$) and suggest that MMFs are just one of the many different classes of filaments discussed in the literature today.  Many MMFs are aligned with the Galactic Plane and may be akin to the dark dust lanes seen in Grand Design Spirals. } ",
        "introduction": " ",
        "conclusions": "\\label{sec:2} Nessie \\citep{jac10} is another such MMF; it too exhibits coherent ($<$ 3.4 km s$^{-1}$) velocity structure over about 80 pc and has a total mass of about 10$^{4}$ \\Msun~as traced by dense gas, comparable to G32.02$+$0.06.  Nessie is also parallel to the Galactic Plane, but slightly offset in latitude from the mid- plane ($|b|$ $\\sim$ -0.4$^{o}$).  There are at least a handful of other MMFs aligned with the Galactic Plane (Tackenberg et al., submitted).  If a majority of MMFs are aligned with the Galactic Plane this indicates that Galactic processes such as shear and spiral density waves are more important than super-bubbles in their formation.  Their alignment with the Galactic Plane is analogous to dark dust lanes along spiral arms seen in face-on galaxies.  G32.02$+$0.06 represents just one example of an MMF in the Galactic mid-plane formed by the compression of previous generations of massive stars.  We classify such filaments (M $\\geq 10^{4}$ \\Msun, length $\\geq$ 10 pc, $\\bigtriangleup$v $\\leq$ 5 km s$^{-1}$) as MMFs and suggest that they represent just one category of the oft-discussed ``filaments\" in the literature of late."
    },
    "1208/1208.5054_arXiv.txt": {
        "abstract": "{% Distances from the tip of the red-giant branch (TRGB) in the halo Population of galaxies  -- calibrated through RR~Lyr stars as well as tied to Hipparcos parallaxes and further supported by stellar models -- are used to determine the luminosity of six nearby type Ia supernovae (SN 2011fe, 2007sr, 1998bu, 1989B, 1972E, and 1937C). The result is $\\langle M^\\mathrm{corr}_{V}\\rangle=-19.41\\pm0.05$. If this value is applied to 62 SNe\\,Ia with $3000< v < 20,000\\kms$ a large-scale value of the Hubble constant follows of $H_{0}=64.0\\pm1.6\\pm2.0$. The SN \\textit{HST} Project gave $H_{0}=62.3\\pm1.3\\pm5.0$ from ten Cepheid-calibrated SNe\\,Ia (Sandage et~al. 2006). The agreement of young Population~I (Cepheids) and old, metal-poor Population~II (TRGB) distance indicators is satisfactory. The combined weighted result is $H_{0}=63.7\\pm2.3$ (i.e.\\ $\\pm3.6\\%$). The result can also be reconciled with WMAP5 data (Reid et~al. 2010). } ",
        "introduction": "\\label{sec:1} The observational determination of the precise large-scale value of the Hubble constant ($H_{0}$) is decisive for understanding the nature of dark energy and for determining the equation of state $\\omega$ and other cosmological parameters \\citep[e.g.][]{Suyu:etal:12}. The most direct route to achieve this goal is through the Hubble diagram of supernovae of type Ia (SNe\\,Ia) which are the tightest standard candles known and trace the linearity of the cosmic expansion out to $z>1$. But for the numerical value of $H_{0}$ the intrinsic luminosity of SNe\\,Ia is also needed. So far their mean luminosity has been derived from a few SNe\\,Ia in nearby Population~I parent galaxies whose distances can be derived from the period-luminosity relation of Cepheids. Present results from calibrated SNe\\,Ia, however, diverge considerably giving, e.g., $H_{0}=71\\pm6$ using only SNe\\,Ia from the \\textit{HST} Project for $H_{0}$ \\citep{Freedman:etal:01}, $62.3\\pm5.2$ from the SN \\textit{HST} Project \\citep{STS:06}, and $73.8\\pm2.4$ \\citep{Riess:etal:11}. The divergence is almost entirely caused by different Cepheid distances: of the ten Cepheid distances used for the calibration of SNe\\,Ia by \\citet{STS:06} six are also in the list of \\citet{Freedman:etal:01} and four (out of 8) are in the list of \\citet{Riess:etal:11}. The distances of the two latter samples are shorter on average by $0.41\\pm0.10$ and $0.32\\pm0.08\\mag$, respectively, than those by \\citet{Saha:etal:06}. The reason for these differences, that perpetuate of course into \\textit{all} Cepheid-calibrated distance indicators, are different interpretations of metallicity effects on the shape and the zero point of the period-color relation and the period-luminosity relation and the related question of internal absorption; these points are further discussed by \\citet[][and references therein]{TR:11}.  The unsatisfactory situation with Cepheid-calibrated SNe\\,Ia calls for an independent calibration. Therefore an alternative SN\\,Ia calibration is pursued here using tip of the red-giant branch (TRGB) distances. The two routes to $H_{0}$ are truly independent because Cepheids belong to the young Population~I whereas the TRGB is a feature of the old Population~II.  Over the last 70 years the TRGB of the \\textit{old, metal-poor} Population has developed into a powerful distance indicator. By now it fulfills the long-standing hope for a large-scale distance scale based on only Population~II objects because the inherent disadvantage of the TRGB, i.e.\\ its limited range, can be overcome with the first SNe\\,Ia with known TRGB distances. Thus the TRGB will provide an independent check on the Population~I distance scale that relies so far almost entirely on Cepheids.  The TRGB has several advantages. It is physically well understood and empirically exceptionally well calibrated. It needs fewer observations than variable stars, the problem of absorption in the halo of the host galaxy is minimal, and within the relevant metallicity range its luminosity is nearly independent of color. The observation of the TRGB in the galaxian halos obliterate the danger of blends, a much debated problem in case of Cepheids. Moreover the determination of the TRGB -- that is defined as a cutoff magnitude -- is immune against selection effects (Malmquist bias), a problem that has severely hampered the application of distance indicators with substantial internal dispersion.  In the following all magnitudes are corrected for Galactic absorption \\citep{Schlegel:etal:98} with $R_{B}=4.1$, $R_{V}=3.1$, and $R_{I}=1.8$. Reduced values of 3.65, 2.65, and 1.35, respectively, are adopted for the absorption of the SNe\\,Ia in their host galaxies as further discussed in Sect.~\\ref{sec:3}. The SN magnitudes refer to the maxima in $B$, $V$, and $I$, respectively, if not otherwise stated. The colors $(B\\!-\\!V)^{\\max}$ are defined as $(B^{\\max} - V^{\\max})$.  The method of the TRGB is discussed in Sect.~\\ref{sec:2}. The TRGB distances and optical luminosities of six SNe\\,Ia are derived in Sect.~\\ref{sec:3}. In Sect.~\\ref{sec:4} the calibrated SN\\,Ia luminosity is applied to distant SNe\\,Ia leading to the value of $H_{0}$. $JHK_{s}$ luminosities are briefly discussed in Sect.~\\ref{sec:5}. The discussion and conclusions follow in Sect.~\\ref{sec:6}. ",
        "conclusions": "\\label{sec:6} SNe\\,Ia play a special role for the determination of the cosmic value of $H_{0}$ because no other (relative) distance indicators trace the cosmic expansion field with so small dispersion well beyond all possible local effect. The next best route is offered by galaxy clusters with multiple $I$-band Tully-Fisher distances, although this method has large intrinsic dispersion and is therefore particularly susceptible to selection bias. The group of R. Giovanelli has determined such distances for 28 clusters with 26 cluster members per cluster on average \\citep{Masters:etal:06}. After a very careful and complex correction for various selection effects they define the slope of the galaxy-type-dependent $m_{I}-\\log(\\rm{line width})$ relation. The ensuing \\textit{relative} cluster distances define the Hubble line with a surprisingly small scatter of only $0.15\\mag$ over a velocity range from 2000 to just $10,000\\kms$. However the three nearest clusters with $v<2000\\kms$ lie $0.4\\mag$ below the Hubble line \\citep[][Fig.~8]{TSR:08a} implying a 20\\% decrease of $H_{0}$ around $2000\\kms$, which is impossible in view of the linearity of the Hubble flow \\citep[e.g.][]{SRT:10}. The discontinuity of $H_{0}$ is a clear indication that the objective definition of cluster samples could not be maintained in nearby clusters. \\citet{Masters:etal:06} have zero-pointed their relative cluster distances by means of 15 nearby field galaxies with Cepheid distances from \\citet{Freedman:etal:01} and of two additional galaxies from other sources. (Their Cepheid moduli carry a special metallicity correction and are larger on average than those given by \\citet{Freedman:etal:01} by $0.10\\mag$; they are still smaller than those given by \\citet{Saha:etal:06} by $0.15\\mag$ on average). In this way they derived a value of $H_{0}=74\\pm2$. However, we strongly object to the calibration procedure: the objectively homogenized cluster sample cannot be compared with the sample of nearby calibrators that are chosen for the sole reason that Cepheid distances and suitable line width data are available for them.  \\citet*{Hislop:Mould:11} have used the data for 15 clusters from \\citet{Masters:etal:06}, but their $I$-band Tully-Fisher relation is calibrated with 13 local field galaxies with available TRGB distances. Six of these calibrators (NGC 247, 3351, 3368, 3621, 3627, and 4826) are questionable because the respective TRGB distances published by \\citet{Mould:Sakai:08,Mould:Sakai:09a,Mould:Sakai:09b} and \\citet{Sakai:etal:04} are larger by $0.34\\mag$ on average. But this is of minor importance compared with the principal objection that -- as in the case of \\citet{Masters:etal:06} -- the samples of calibrators and cluster galaxies constitute entirely differently selected samples. If anything, the situation is here even worse, because \\citeauthor{Hislop:Mould:11} use only ~30\\% of the available cluster members thereby questioning the corrections for selection bias derived by \\citet{Masters:etal:06}. For these reasons the value of $H_{0}=79\\pm2$ of \\citet{Hislop:Mould:11}  should be given low weight.  To support the present solution it is noted that for the six galaxies in Table~\\ref{tab:01} also Cepheid distances are available. The corresponding data are set out in the self-explanatory Table~\\ref{tab:03}. The resulting mean Cepheid-calibrated luminosity of the six SNe\\,Ia is $M_{V}^\\mathrm{corr}=-19.40\\pm0.06$ with a dispersion of $0.15\\mag$ (i.e.\\ larger than in case of the TRGB calibration). Hence the calibration through the TRGB distances and the adopted Cepheid distances yields -- in fortuitous agreement -- the same mean luminosity of SNe\\,Ia to within $0.01\\mag$.   \\begin{table} \\begin{center} \\caption{The luminosity of six calibrating SNe\\,Ia from Cepheid distances} \\label{tab:03} \\scriptsize \\begin{tabular}{llrccc} \\hline \\hline \\noalign{\\smallskip} \\multicolumn{1}{c}{SN}                          & \\multicolumn{1}{c}{Gal.}                        & \\multicolumn{1}{c}{$m^{\\rm corr}_{V}$}          & \\multicolumn{1}{c}{$(m\\!-\\!M)_{\\rm Ceph}$}      & \\multicolumn{1}{c}{Ref.}                        & \\multicolumn{1}{c}{$M^{\\rm corr}_{V}$}          \\\\ \\multicolumn{1}{c}{(1)}                         & \\multicolumn{1}{c}{(2)}                         & \\multicolumn{1}{c}{(3)}                         & \\multicolumn{1}{c}{(4)}                         & \\multicolumn{1}{c}{(5)}                         & \\multicolumn{1}{c}{(6)}                         \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 2011fe & N5457   & $ 9.93\\,(09)$ & $29.28\\,(08)$ & 1 & $-19.35\\,(11)$ \\\\ 2007sr & N4038   & $12.28\\,(17)$ & $31.66\\,(08)$ & 2 & $-19.38\\,(19)$ \\\\ 1998bu & N3368   & $11.01\\,(12)$ & $30.34\\,(11)$ & 3 & $-19.33\\,(16)$ \\\\ 1989B  & N3627   & $10.94\\,(11)$ & $30.50\\,(09)$ & 3 & $-19.56\\,(14)$ \\\\ 1972E  & N5253   & $ 8.37\\,(11)$ & $28.05\\,(27)$ & 3 & $-19.68\\,(29)$ \\\\ 1937C  & I4182   & $ 8.92\\,(16)$ & $28.21\\,(09)$ & 3 & $-19.29\\,(18)$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\multicolumn{5}{l}{straight mean}             & $-19.43\\pm0.06$ \\\\ \\multicolumn{5}{l}{weighted mean}             & $-19.40\\pm0.06$ \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{center} \\tablefoot{% (1) \\citealt{TR:11}; (2) \\citealt{Riess:etal:11} from $H$ magnitudes; (3) \\citealt{Saha:etal:06}. } \\end{table}   It is also noted that the TRGB distances of 78 field galaxies with $v_{220}>280\\kms$ give a local value of $H_{0}=62.9\\pm1.6$ (\\citeauthor*{TSR:08b}). A slightly augmented sample of \\citet{Saha:etal:06}, comprising now 29 Cepheid distances with $280<v_{220}<1600\\kms$, yields $63.4\\pm1.8$ (\\citeauthor*{TSR:08b}). The large-scale value from the summary paper of the \\textit{HST} Supernova Project, based on ten SNe\\,Ia with Cepheid distances and a total of 62 standardized SNe\\,Ia with $2000<v<20,000\\kms$, that resulted in $H_{0}=62.3\\pm1.3\\pm5.0$ \\citep{STS:06}, has been mentioned before.  Additional support for a rather low value of $H_{0}$ comes from \\citet{Reid:etal:10} who have derived -- on the assumption of a $\\Lambda$CDM model, without using $H_{0}$ as a prior -- a value of $H_{0}=65.6\\pm2.5$ by combining the 5-year WMAP data \\citep{Dunkley:etal:09} with the Hubble diagram of the SN\\,Ia Union Sample \\citep{Kowalski:etal:08} and with a sample of over 110,000 luminous red galaxies (LRG) from the Sloan Digital Sky Survey DR7.  A new prospect to \\textit{directly} measure $H_{0}$ at cosmological relevant distances comes from water maser distances of galaxies beyond $\\sim\\!3000\\kms$; the first two galaxies, UGC\\,3789 and NGC\\,6264 yield a  mean value of $H_{0}=67$ with a still large error of $\\pm6$ \\citep{Braatz:etal:10,Kuo:11,Henkel:etal:12}. Further progress along this novel route is to be expected. -- Present results from strongly lensed quasars \\citep[e.g.][]{Vuissoz:etal:08} and from the Sunyaev-Zeldovich effect \\citep[][]{Holanda:etal:12} give even larger error margins on $H_{0}$.  The main conclusion is that the TRGB distances of six SNe\\,Ia require a large-scale value of the Hubble constant of $H_{0}=64.0\\pm2.6$ (including random and systematic errors). Suggested values as high as $H_{0}>70$ are unlikely in the light of the TRGB. If the present result is combined with the independent value of $H_{0}= 62.3\\pm5.2$ from ten Cepheid-calibrated SNe\\,Ia \\citep{STS:06} one obtains a weighted mean of $H_{0}=63.7\\pm2.3$, which is our best estimate. The result is in good agreement with local determinations of $H_{0}$ and the linearity of the expansion field \\citep{SRT:10}.  The present result will be further improved by future TRGB distances of galaxies that have produced a standard SN\\,Ia. It is desirable that the search for the TRGB in NGC\\,3368 and 3627 in the Leo~I group will be repeated in some uncontaminated halo fields in order to definitely proof their group membership. The next easiest targets at present are the Branch normal SN\\,Ia in the Virgo cluster galaxies NGC\\,4419, 4501 (with the highly absorbed SN\\,1999cl), 4526, and 4639 (yet on the far side of Virgo); the two rather old SNe\\,Ia in NGC\\,4496A and 4536 in the W-cloud are less attractive. Somewhat more difficult are the Fornax cluster galaxies NGC\\,1316 with five SNe\\,Ia of which three are Branch-normal \\citep[1980N, 1981D, and 2006dd;][]{Stritzinger:etal:10}, and the galaxies NGC\\,1380 and NGC\\,1448 with one SN\\,Ia each; SN\\,2007on in NGC\\,1404 is as a fast decliner \\citep{Kattner:etal:12} unreliable for calibration purposes. The Virgo and Fornax galaxies will probably be easier and more difficult, respectively, by a few $0.1\\mag$ than NGC\\,4038 where \\citet{Schweizer:etal:08} have demonstrated that the TRGB can be reached.  It is foreseeable that the route to a precision value of $H_{0}$ through TRGB-calibrated SNe\\,Ia will become highly competitive with any Cepheid-based distance scale."
    },
    "1208/1208.0447_arXiv.txt": {
        "abstract": "In this paper we present the various concepts behind the Astro-WISE Information System. The concepts form a blueprint for  general scientific information systems (WISE) which can satisfy a wide and challenging range of  requirements for the data dissemination, storage and processing for various fields in science. We review the main features of the information system and its practical implementation. ",
        "introduction": "\\label{intro}  Digital astronomical catalogues have been built from the very first moment information technology enabled this, e.g. the first Abell catalogue of clusters of galaxies was digitally prepared and printed in the days minus signs were not available in print~\\cite{Abell}. As soon as digital scanning devices became available photographic material was scanned and first large image surveys were digitally published, such as ESO-LV~\\cite{ESO-LV} and the  digitized sky surveys like  DPOSS~\\cite{DPOSS}, succeeded by the Palomar-Quest Survey~\\cite{Quest}, in the 90's followed by CCD-based surveys  such as the 2MASS survey~\\cite{2MASS},  the Sloan Digital Sky Survey SDSS~\\cite{SDSS}.  The rapid accumulation of astronomical digital data and its public dissemination was, compared to other disciplines, achieved at  an early stage, thanks to an open and collaborative astronomical world community who adopted the FITS image data format as early as 1979~\\cite{FITS}.  The CDS, initially Centre de Donn\\'ees Stellaires and later renamed into Centre de Donn\\'ees astronomiques de Strasbourg took the lead in Europe to collect and disseminate the ever growing data sets of a zoo of astronomical observatories and projects. VizieR webservices nowadays provide access to over 9000 catalogues. Numerous astronomical data centers followed,  developing specialized services, for example, the Infrared Processing and Analysis Center and the bibliographical SAO/NASA Astrophysics Data System.   In the early 2000's it was realized that the ever growing data volumes require new approaches: the community becomes the data provider and the International Virtual Observatory and its European branch, the European  Virtual Observatory, developed standards and interfaces to allow individual data centers to publish their catalogs and images in a common framework providing worldwide access to users who can query, cross-match and visualize multiple databases via the internet. A highlight formed the overplotting of data from different experiments, like X-ray satellites and optical ground-based observatories  with some keystrokes on the Aladin interactive software sky atlas.   While the Virtual Observatory focused on the dissemination of published data it was also realized in the early 2000's that the upcoming data deluge required a new approach to the handling of the data stream from the telescope to the science-ready result.  Modern experiments not only involve more data, but also require increasing precision on the various calibrations. The dependency of the final result (sometimes only a few numbers, like cosmological parameter values)  on time-variable  calibrations of very large data sets,   involving the evaluation by  large research teams distributed in smaller groups over various sites sets the basic requirements to the system handling the data. While in classical systems data is delivered in various releases to the public, often re-processing the whole set with higher versions of the code, the  high data rates of modern experiments require an alternative approach, where the up-to-date result is derived on demand by the user.   Thus, we set out to design and implement an integrated {\\it datacentric}  system Astro-WISE  in which the processing, storage and administration is integrated in a single environment, providing a {\\it living} system to both the data producers and the customers. Early reports of this development and implementation have been given in~\\cite{V1} and~\\cite{V2}. To reach this goal,  traceability of each individual data item handled by pipelines or any piece of code is carefully maintained, every data item beyond pixel value is  kept as  Metadata and made persistent and distributed in a relational database with an object-oriented view, mapping all the dependencies. ",
        "conclusions": "The Astro-WISE information system, the first information system in Astronomy, proved to be a reliable and flexible tool for the data processing. Originally developed to process the data of KIlo Degree Survey (KIDS\\footnote{http://www.astro-wise.org/projects/KIDS/}) it triggered development of the unique approach to the architecture of scientific information systems (WISE approach).   Both Astro-WISE and the WISE approach are living systems which are open to improvements. For the last 2 years further development of the WISE approach is hosted by Target Holding\\footnote{http://www.rug.nl/target}. Target Holding is an expertise center in the Northern Netherlands which is building a cluster of sensor network information systems and provides cooperation between a number of scientific projects and business partners like IBM and Oracle. Target creates and supports a hardware infrastructure for hosting tens of Petabytes of data for projects in astronomy, medicine, artificial intellegence and biology.  In this development the WISE approach was used to create new information systems extending the original Astro-WISE on new data models, new data storage and processing capacities and new fields."
    },
    "1208/1208.2990_arXiv.txt": {
        "abstract": "By combining large-scale mosaics of ROSAT PSPC, XMM-Newton, and Suzaku X-ray observations, we present evidence for large-scale motions in the intracluster medium of the nearby, X-ray bright Perseus Cluster. These motions are suggested by several alternating and interleaved X-ray bright, low-temperature, low-entropy arcs located along the east-west axis, at radii ranging from $\\sim10$ kpc to over a Mpc. Thermodynamic features qualitatively similar to these have previously been observed in the centers of cool core clusters, and were successfully modeled as a consequence of the gas sloshing/swirling motions induced by minor mergers. Our observations indicate that such sloshing/swirling can extend out to larger radii than previously thought, on scales approaching the virial radius. ",
        "introduction": "Among the first important {\\sl Chandra} results on galaxy clusters was the surprising discovery of cold fronts - remarkably sharp surface brightness edges, with relatively dense, cool gas on the inner (bright) side and low density hot gas on the outer (faint) side \\citep{Markevitch2000,Markevitch01}. Cold fronts are seen both in merging systems, where the edge appears at the projected location of the density discontinuity separating the hotter intra-cluster medium (ICM) from the low entropy core of an infalling subcluster, and in the centers of cool-core clusters (at $r=$10--400~kpc) with little or no clear signs of recent mergers. In the latter case, multiple arcs, curved around alternating sides of the central gas density peak along a spiral-like pattern, are often seen \\citep[][and references therein]{markevitch2007}.  Cold fronts in cooling cores are believed to be due to the sloshing (or, more accurately, swirling) of the gas in the gravitational potentials of clusters \\citep{Markevitch01}. Numerical simulations show that sloshing can be induced easily by a minor merger, where a subcluster falls in with a nonzero impact parameter \\citep{Tittley05, Ascasibar06}. The gas from the subcluster is typically stripped early during the infall, but its dark matter continues to fall in and, due to the gravitational disturbance, the density peak of the main cluster swings on a spiral-like trajectory relative to the center of mass. As the central ICM is displaced, concentric cold fronts are created where cooler and denser parcels of gas from the center come into contact with the hotter outskirts. The cold fronts associated with sloshing/swirling can persist for Gyrs. It has been proposed that the spiral pattern of cold fronts in the innermost regions of cool core clusters signals the presence of large-scale bulk spiral flows \\citep{Keshet2011}.  The ICM in the core of the Perseus Cluster of galaxies (A426), which is the brightest, extended extragalactic X-ray source, also shows such a spiral pattern \\citep{Fabian00,Churazov00,Churazov03}. A deep 1.4 Ms {\\sl Chandra} observation has confirmed the presence of a cold front located about 100~kpc west of the cluster core \\citep{Fabian2011}. Previous Einstein Observatory and {\\it ROSAT PSPC} imaging had furthermore suggested the presence of an east-west shift in the centroid of the X-ray isophotes for this system on larger scales of several hundred kpc, which was interpreted as a sign of an ongoing or past merger \\citep{BranduardiRaymont81,Schwarz92,Allen92}. In this Letter, we use both archival ROSAT PSPC and XMM-Newton data, as well as a recently obtained Suzaku mosaic of the Perseus Cluster to investigate the origin of these asymmetries in unprecedented detail, from the cluster core out to large radii. We examine possible links between the smaller-scale spiral pattern previously seen in the core and new features reported here on much larger scales, and discuss the implications for the past and current dynamical state of the cluster, and the physical properties of the ICM.  We assume a $\\Lambda$CDM cosmology with $\\Omega_m$=0.27, $\\Omega_\\Lambda$=0.73, and H$_0$=70 km/s/Mpc. At the redshift of the cluster, z=0.0179 \\citep{Struble99}, one arcminute corresponds to 21.54 kpc. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=1.05\\columnwidth]{sb_t_s_avg.pdf} \\caption{Surface brightness, projected temperature, and deprojected entropy profiles obtained from our Suzaku data. The surface brightness profiles are shown individually for the E, SE, and NE directions all sampling the E edge at $\\sim700$~kpc, as well as for the W and SW arms showing the large-scale excess beyond this radius. The temperature and entropy profiles were obtained by averaging the three arms towards the E and two towards the W. For reference, in the entropy panel we overplot the best-fit power-law with an index 1.1, fitted to the data from the remaining three relatively more relaxed arms (S,N,NW). \\label{suzaku}} \\end{figure}  As mentioned above, thermodynamic features qualitatively similar to those presented in this work have been modeled successfully in the centers of cool core clusters. Numerical simulations by \\cite{Tittley05} and \\cite{Ascasibar06} predict that gas sloshing induced by a disturbance of the cluster's gravitational potential produces precisely this type of opposite-and-staggered features, with alternating regions of high X-ray brightness and low temperature on either side of the main cluster's X-ray peak.  We note that, based on the presence of the eastern surface brightness edge alone, we would not be able to distinguish whether the eastern excess is caused by a large-scale displacement of the main cluster gas, or whether it originated from the stripped ICM of an infalling subcluster (this was the hypothesis put forth by \\citealt{Schwarz92} when first describing this feature). However, given the connection of this feature to both smaller and larger scales as described above, we argue that the large-scale sloshing/swirling interpretation is more likely. The eastern excess both begins - in radius - just beyond the edge of the inner western cold front and ends within the concave inner edge of the outer western surface brightness enhancement. The temperature profiles anticorrelate with the X-ray brightness and gas density, as predicted by numerical models of gas sloshing in cluster cool cores.  Gas motions in the core of the Perseus Cluster (inner 100~kpc) have already been indicated by the lack of resonant scattering in the 6.7~keV He-like iron line \\citep{churazov2004}, which implies a range in velocities of at least half of the sound speed. Significant swirling motion in the ICM may also affect the azimuthal distribution of the different generations of bubbles of relativistic plasma inflated by the central AGN. Future observations with the high resolution X-ray spectrometers on the upcoming Japanese-American Astro-H satellite will allow us to accurately map the line-of-sight component of these gas motions in the cluster core.  The challenge of the results presented here is to understand how gas sloshing/swirling can reach well beyond the cluster's cool core. Already the cold front at $700$~kpc is well outside the cooling radius (approximately 200~kpc), and thus outside the radial range where the temperature dips in the cluster center. The western surface brightness excess reaches even further in radius, out to well beyond 1~Mpc. Numerical simulations have so far only attempted to reproduce sloshing cold fronts within the cool-core region characterized by a steep density gradient and a decrease of the temperature at small radii. Based on these models, \\citet{Ascasibar06} conclude that ``with the possible exception of short moments of the sub-cluster flyby generating a conical wake, the cluster stays very symmetric on large scales; the only structure is edges in the center.''  \\citet{Roediger11} performed a quantitative modeling of gas sloshing in the core of the Virgo Cluster, which exhibits spiral-like features in the surface brightness and gas temperature distributions which are qualitatively similar to those presented here, although on much smaller spatial scales. They found that typical signatures of the oscillation of the main cluster gas do include a large-scale brightness asymmetry. It is interesting to note that this brightness excess, which was also associated with a lower average gas temperature compared to other azimuths, did reach beyond the cooling radius of the Virgo Cluster according to the numerical simulations, out to as far as 400~kpc (see Fig. \\ref{sim}). However, the outermost surface brightness edge that is observed in the Virgo Cluster is only 90~kpc away from its center \\citep{Simionescu10}, compared to the location of our eastern edge at $700$~kpc. Moreover, based on Fig. \\ref{sim}, it would appear that in order to reproduce the large-scale east-west asymmetries, the preferred trajectory of the infalling subcluster would be oriented roughly north-south; while this cannot be ruled out, it seems much more likely in the case of the Perseus Cluster for a perturber to fall in along the east-west direction, which coincides with the major axis of the cluster, as well as with an elongated chain in the galaxy distribution on even larger scales, as part of the Perseus-Pisces supercluster \\citep{Haynes86}.  It will therefore be important to perform numerical simulations targeted at reproducing the features associated with the gas sloshing on very large scales seen in the Perseus Cluster, in order to understand whether this can be explained by an area of the merger parameter space that has not been modeled before, or whether more complex microphysics such as an enhanced gas viscosity is required for these features to be replicated. Such numerical simulations will also allow us to determine whether the asymmetries seen in Perseus could have been caused by a recent merger, or whether they are due to longer-lived modes of oscillation trapped in the gravitational potential of the cluster. Based on cosmological simulations, the presence of significant turbulence and large-scale bulk motions in the ICM, especially at large radii, is not surprising \\citep{Burns10,Nagai07,Rasia06}, but the apparent {\\sl coherence} of the large-scale spiral-like structure ranging over at least two decades in radius is unexpected.  The Perseus Cluster has a pronounced cool core, with a sharp peak in X-ray surface brightness, a decreasing temperature towards the center, and obvious signs of ongoing AGN feedback processes \\citep[e.g.][]{Boehringer93,Fabian00}. The spiral-shaped brightness enhancement seen in the very center on scales of 10--100~kpc seems to be connected to the larger-scale sloshing, which shows that the motions induced by the merger penetrate the cool core. Therefore, if a relatively powerful merger is required to have induced the large-scale sloshing discussed here, it will be important to understand the circumstances under which this merger does not destroy the cluster's cool core. This analysis may provide more general clues towards the survivability of cluster cool cores during mergers \\citep[for a discussion, see][]{Burns08,Million10}."
    },
    "1208/1208.0392_arXiv.txt": {
        "abstract": "Recent astronomical data strongly suggest that a significant part of the dark matter, composing the Local Group and Virgo Supercluster, is not incorporated into the galaxy haloes and forms diffuse components of these galaxy clusters. Apparently, a portion of the particles from these components may penetrate into the Milky Way and make an extragalactic contribution to the total dark matter containment of our Galaxy.  We find that the particles of the diffuse component of the Local Group are apt to contribute $\\sim 12\\%$ to the total dark matter density near the Earth. The particles of the extragalactic dark matter stand out because of their high speed ($\\sim 600$~{km/s}), i.e. they are much faster than the galactic dark matter. In addition, their speed distribution is very narrow ($\\sim 20$~{km/s}). The particles have isotropic velocity distribution (perhaps, in contrast to the galactic dark matter). The extragalactic dark matter should give a significant contribution to the direct detection signal. If the detector is sensitive only to the fast particles ($v<450$~{km/s}), the signal may even dominate.  The density of other possible types of the extragalactic dark matter (for instance, of the diffuse component of the Virgo Supercluster) should be relatively small and comparable with the average dark matter density of the Universe. However, these particles can generate anomaly high energy collisions in direct dark matter detectors. ",
        "introduction": "It is widely believed that all the dark matter particles (hereafter DMPs), which a terrestrial observer can detect, belong to the Milky Way Galaxy. The main aim of this letter is to dispute this assertion and to show that a remarkable fraction of dark matter particles detected on the Earth does not probably belong to our Galaxy. Although their density is relatively small, as compared with the total dark matter density, their impact into the direct detection signal may even dominate because of high speeds of the particles.  According to the modern cosmological notion, the haloes of giant galaxies, like Milky Way or Andromeda, are regions of local dark matter overdensity, rather than isolated islands. Indeed, the Local Group, along with the haloes of large and dwarf galaxies, contains a significant fraction of dark matter that is not bound in the galaxies and presumably forms a large envelope of the Local Group \\citep{bt}. A significant part of the dark matter of the Virgo Supercluster is also not localized in haloes and probably distributed more or less homogeneously over all the volume of the Supercluster \\citep{karach11}. Some part of this diffuse dark matter (preeminently from the Local Group envelope) penetrates into the central region of our Galaxy and gives a contribution to the direct detection signal, which can even dominate under certain conditions: as we will see, the density fraction of the extragalactic dark matter is relatively small ($\\sim 12\\%$), however its particles should have extremely high speeds, close to the escape velocity ($\\sim 600$~{km/s}) or even higher. It sets off the extragalactic particles from the halo DMPs with much lower average speed. The direct detection signal produced by this component should also have some other characteristic features that will be discussed below. ",
        "conclusions": "Not much is known about the dark matter distribution in the envelope. On the other hand, as we can see from equation (\\ref{14a19}), the result is not strongly dependent on the choice of $i$ and $j$ (as a consequence of the relative smallness of ratio $r_{out}/r_{in}$). It seems reasonable to choose $i$ and $j$ by analogy with the well-known isothermal halo solution ($dM/dr={\\it const}$ and the Maxwell DMP velocity distribution with a temperature, constant over the halo), which corresponds to $i=1$, $j=0$. Substituting these values in combination with $T(r_{in})$, $\\alpha(r_{in})$, and $V$ to (\\ref{14a19}), we obtain the density of the extragalactic dark matter near the Earth $\\rho= 3.7\\times 10^{-2}$~{GeV/cm$^3$}. The speed distribution is notably narrow: the absolute values of all particles fall within $\\Delta V\\simeq 16$~{km/s}. As we have already mentioned, this feature is a consequence of the smallness of $|\\phi(r_{out})-\\phi(r_{in})|$ as compared with $-\\phi(\\lo)$. Therefore, two properties of the velocity distribution are model-independent: the speeds of extragalactic DMPs from the envelope lie in a narrow range, and their angular distribution is isotropic.  The density of the extragalactic dark matter turns out to be fairly high: $3.7\\times 10^{-2}$~{GeV/cm$^3$} is more than $12$\\% of the total dark matter density near the Earth $\\simeq 0.3$~{GeV/cm$^3$} \\citep{gorbrub1}. This brings up a question: How reliable is the estimation? Above we have already discussed the approximation of the system by a spherically symmetric model and found it acceptable. The premise that $f(r_0)$ terminates abruptly at $r_{in}$ and $r_{out}$ is also unphysical. Undoubtedly, our result is assessed; however, it cannot be called optimistic. Indeed, as the dimensional method shows, for any envelope model the density of the extragalactic dark matter is, with an accuracy of a numerical factor, equal to \\begin{equation} \\label{14a21} \\rho\\propto \\frac{M_{env} v_{esc}}{\\langle\\alpha\\rangle^2 \\langle T\\rangle} \\end{equation} where $\\langle\\alpha\\rangle$ and $\\langle T\\rangle$ are the average values of the respective quantities. Our calculations confirm this dependence: it can be easily obtained from (\\ref{14a13}). $v_{esc}$ is almost independent on the model choice. $\\langle T\\rangle$ is essentially defined by the size of the Milky Way Roche lobe, and thus is also more or less model-independent. The main source of the uncertainty is envelope mass $M_{env}$. We proceed from the assumption of \\citet{collision} that $M_{env}$ is approximately equal to the masses of the galaxies of the Local Group. We used the highest possible value (\\ref{14a7}) for $\\alpha$: if $\\alpha$ was higher, a significant part of the envelope would rapidly evaporate. Since $\\rho\\propto\\alpha^{-2}$, this choice is conservative.  Thus there are two possible situations, when (\\ref{14a19}) significantly overestimates the density of the extragalactic dark matter. It may be so, if the envelope mass is in fact much lower than the masses of the Local Group member galaxies. The strong overestimation may also appear, if the angular momentum distribution of the envelope DMPs differs greatly from the Gaussian (\\ref{14a1}), i.e., almost all the particles have circular orbits. Such a supposition seems highly improbable. First of all, it is in sharp contrast to N-body simulation results \\citep{mo09}. There is also a good indirect counterargument: the largest Local Group member M31 has quite low angular momentum and, consequently, very oblong orbit \\citep{1959ApJ...130..705K}. It is plausible that Milky Way and M31 will finally experience a central collision. Thus the presence in the diffuse component of the Local Group of a bulk of dark matter particles that have very oblong orbits and can reach the Earth seems quite possible.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=0]{14fig2.eps}} \\caption{The ratio of the direct detection signal produced by the mixture $\\sim 12.3\\%$ of extragalactic component and $87.7\\%$ of the galactic DMPs to the signal produced by pure galactic dark matter. The fraction and velocity distribution of the extragalactic component were calculated in accordance with (\\ref{14a19}). We considered two models of the velocity distribution of the galactic DMPs: Maxwell (\\ref{14a23}) (solid line) and anisotropic (\\ref{14a24}) (dashed line). The extragalactic dark matter almost does not affect the signal, if $\\upsilon_{min}<300$~{km/s}, but totally dominates above $450-500$~{km/s}.} \\label{fig2} \\end{figure}   The $12\\%$ extragalactic component of the total dark matter density can be especially important for the direct dark matter search. The direct search is based on the detection of the collisions of dark matter particles with nuclei of the target. The signal is sensitive to the velocity distribution: roughly speaking \\cite{Belanger}, it is proportional to \\begin{equation} I(\\upsilon_{min})=\\int^\\infty_{\\upsilon_{min}}\\frac{\\tilde n(\\upsilon)}{\\upsilon}\\; d^3 \\vec\\upsilon \\label{11b8} \\end{equation} Here  $\\tilde n(\\upsilon)$ is the distribution in the Earth's frame of reference: it should be obtained from (\\ref{14a20}), (\\ref{14a23}), or (\\ref{14a24}) by a Galilean transformation. $\\upsilon_{min}$ is the minimal DMP speed, to which the detector is sensitive (see details in \\citet{Belanger}). \\begin{equation} \\label{14a22} \\upsilon^2_{min}\\simeq \\frac{E_A}{2} \\frac{(m_\\chi+m_A)^2}{m_A m^2_\\chi} \\end{equation} where $m_\\chi$ and $m_A$ are the DMP and the detector nucleus mass respectively, $E_A$ is the detector activation energy, depending on its construction. In order to estimate the influence of the extragalactic dark matter to the direct detection signal, we should define a model for the velocity distribution of the galactic component. The Maxwell distribution is now routinely used, mainly because of its simplicity: \\begin{equation} n(u) = \\frac{1}{(\\sqrt{\\pi} \\upsilon_{\\odot})^3} \\exp \\left(- \\frac{u^2}{\\upsilon_{\\odot}^2}\\right) \\label{14a23} \\end{equation} $\\upsilon_{\\odot}$ is the orbital speed of the Solar System. There are strong reasons to suppose, however, that the distribution of the galactic DMPs is strongly anisotropic and looks like \\begin{equation} n(u)= \\frac{\\exp\\left(-\\dfrac{u^2_\\tau}{2\\sigma^2_0}\\right)}{2\\pi^2\\sigma^2_0\\sqrt{u^2_{max}-u_r^2}} \\label{14a24} \\end{equation} where $u_r\\in [-\\upsilon_{max};\\upsilon_{max}]$, $u_{max}\\simeq 560$~{km/s}, $\\sigma_0=80$~{km/s}, $u_r$ and $u_\\tau$ are radial and tangential components of the particle velocity, respectively \\citep{2011MNRAS.417L..83B}.  We calculated the signal produced by mixture of $3.7\\times 10^{-2}$~{GeV/cm$^3$} of the extragalactic dark matter and $0.263$~{GeV/cm$^3$} of the galactic one and divided it by the signal produced by the pure galactic dark matter with the same total density ($0.3$~{GeV/cm$^3$}). We used both models of the galactic DMP distribution: Maxwell (\\ref{14a23}) and anisotropic (\\ref{14a24}). The two ratios are represented in Fig.~\\ref{fig2} by the solid line (Maxwell velocity distribution of the galactic dark matter particles) and by the dashed line (anisotropic velocity distribution). One can see that the signal is scarcely affected by the presence of the extragalactic component, if $\\upsilon_{max}<300$~{km/s}. However, the situation drastically changes for higher $\\upsilon_{max}$; if $\\upsilon_{max}$ is larger, than $450-500$~{km/s}, the extragalactic signal dominates. This is not particularly surprising: all the extragalactic particles are faster than $600$~{km/s}, while the number of the galactic particles rapidly drops above $\\sim 450$~{km/s}. Hence the impact of the extragalactic dark matter to the direct detection signal can be very important, especially if the DMP mass $m_{\\chi}$ is small. Indeed, if $m_{\\chi}$ is small, $\\upsilon_{max}$ is high (\\ref{14a22}), i.e., we can detect only the fastest DMPs. For instance, DAMA collaboration reports \\citep{dama} about the detection of a signal produced by $\\sim 10$~GeV weakly interacting massive particles. We shall not discuss here the question of the nature of the signal (other detectors do not confirm the result \\citep{xenon}). It should be recorded, however, that $\\upsilon_{max}>450$~{km/s} for the majority of detectors, if the DMP is so light, i.e., the extragalactic component signal should totally dominate.  Notice that we should, strictly speaking, have cut distributions (\\ref{14a23}) and (\\ref{14a24}) at $u=v_{esc}$. However, the fraction of the particles with $u>v_{esc}$ is negligible for both the distributions, and the cutting would hardly affect the result; the impact of the extragalactic component would be even slightly higher.  In conclusion, let us briefly consider the extragalactic dark matter that does not belong to the Local Group. As recent astronomical observations imply \\citep{karach11}, the dark matter of the Virgo Supercluster, in addition to galaxies and their groups, forms a large diffuse component. We do not know its velocity and space distributions, but it seems reasonable to assume that the dark matter is distributed more or less uniformly, and the velocity dispersion of the DMPs is comparable with that of the observable members of the Supercluster ($v_{\\infty}\\sim 500$~{km/s}). The measurements estimate the average density of the diffuse component as $\\rho\\sim 10^{-6}$~{GeV/cm$^3$}. The gravitational field of the Local Group should increase this quantity near the Earth. However, we can roughly estimate the enhancement as $1+v^2_{esc}/v^2_{\\infty}$, where $v_{esc}\\simeq 650$~{km/s} is the escape velocity from the Solar System orbit \\citep{2012MNRAS.420..590B}. Thus the density of the Supercluster dark matter is approximately $3$ times higher near the Solar System, but yet hardly exceeds $10^{-5}$~{GeV/cm$^3$}. This value is so low, that it may scarcely be of interest for modern experiments. On the other hand, the Supercluster DMPs are particularly energetic ($v>1000$~{km/s}) and hence may give a very characteristic signal.  To summarize:  1) The particles of the diffuse component of the Local Group are apt to contribute $\\gtrsim 10\\%$ to the total dark matter density near the Earth.  2) The particle speeds are $\\sim 600$~{km/s}, i.e. they are much faster than the galactic DMPs. The particles have isotropic velocity distribution (perhaps, in contrast to the galactic dark matter); their speed distribution is very narrow ($\\Delta V\\sim 20$~{km/s}).  3) The extragalactic dark matter should give a significant contribution to the direct detection signal. If the detector is sensitive only to the fast particles ($v> 450$~{km/s}), the signal may even dominate.  4) The density of other types of the extragalactic dark matter (for instance, of the DMPs forming the diffuse component of the Virgo Supercluster) should be relatively small and comparable with the average dark matter density of the Universe. However, these particles can generate anomaly high-energy collisions in direct dark matter detectors.  Financial support by Bundesministerium f\\\"ur Bildung und Forschung through DESY-PT, grant 05A11IPA, is gratefully acknowledged. BMBF assumes no responsibility for the contents of this publication."
    },
    "1208/1208.5975_arXiv.txt": {
        "abstract": "We discuss the detectability of gravitational waves with a time dependent mass contribution, by means of the stochastic gravitational wave observations. Such a mass term typically arises in the cosmological solutions of massive gravity theories. We conduct the analysis based on a general quadratic action, and thus the results apply universally to any massive gravity theories in which modification of general relativity appears primarily in the tensor modes. The primary manifestation of the modification in the gravitational wave spectrum is a sharp peak. The position and height of the peak carry information on the present value of the mass term, as well as the duration of the inflationary stage. We also discuss the detectability of such a gravitational wave signal using the future-planned gravitational wave observatories. ",
        "introduction": "Since the pioneering model of massive gravity was proposed by Fierz and Pauli \\cite{Fierz:1939ix}, numerous attempts have been made to introduce a non-zero mass to graviton. This issue has been attracting a great deal of interest, partly because graviton mass may provide an alternative explanation for the acceleration of our universe. Namely, instead of attributing it to the existence of dark energy whose origin is still unknown, the acceleration may be simply due to modification of gravity.   Until recently, however, it has been thought that such modification by non-zero graviton mass is extremely difficult, typically leading to the emergence of ghost degrees of freedom and various pathologies. As first shown by Boulware and Deser~\\cite{Boulware:1973my}, the source of these difficulties is associated with a helicity-$0$ mode in the gravity sector, which is absent at the quadratic order of the Fierz-Pauli massive gravity but revives at the higher order and acts as a ghost, dubbed the {\\it Boulware-Deser (BD) ghost}. A remedy for this difficulty was recently proposed by Refs.~\\cite{deRham:2010ik,deRham:2010kj}, in which the would-be BD ghost is eliminated by adding higher-order correction terms order by order and by resumming them into a three parameter expression. Non-existence of the BD ghost in this theory was discussed and proved in Refs.~\\cite{Hassan:2011hr,deRham:2011rn,deRham:2011qq,Hassan:2011tf,Hassan:2011ea,Mirbabayi:2011aa,Golovnev:2011aa, Kluson}. Various solutions in the theory were studied in Refs.~\\cite{deRham:2010tw, Koyama:2011xz, Nieuwenhuizen:2011sq, Koyama:2011yg, Chamseddine:2011bu, D'Amico:2011jj, Gumrukcuoglu:2011ew, Koyama:2011wx, Comelli:2011wq, Volkov:2011an, vonStrauss:2011mq, Comelli:2011zm, Berezhiani:2011xx,Brihaye:2011aa,Kobayashi:2012fz}. Therefore this nonlinear theory of massive gravity can be considered as a nonlinear completion of the Fierz-Pauli theory, which is the simplest and the oldest linear massive gravity theory. Having this in mind, we consider it quite natural to suppose that other massive gravity theories or, more generally, Higgs phases of gravity might find their nonlinear completions in the future.   Such examples include, but are not restricted to, ghost condensate~\\cite{ArkaniHamed:2003uy,ArkaniHamed:2005gu} and Lorentz-violating massive gravity theories~\\cite{Rubakov:2004eb,Dubovsky:2004sg}. In these examples, propagating degrees of freedom do not form a representation of $4$-dimensional Poincar\\'e symmetry even in the exact Minkowski background. Instead, they form representations of $3$-dimensional rotational and translational symmetries. As a result, physical degrees of freedom are classified into scalar, vector and tensor parts, according to their transformation properties under the $3$-dimensional spatial rotation. Modes in different classes may behave rather differently. For instance, in the case of ghost condensate in Minkowski background, propagating degrees of freedom in the gravity sector are two tensor modes and one scalar mode. The two tensor modes behave exactly like those in general relativity (GR) while the scalar mode has an unusual dispersion relation, i.e.\\ non-relativistic dispersion relation without a mass gap. On the other hand, in a particular class of models among those proposed in \\cite{Dubovsky:2004sg}, propagating degrees of freedom in the gravity sector are two tensor modes only, and their dispersion relation is massive and relativistic. From the point of view of the symmetry, i.e.\\ $3$-dimensional spatial rotational symmetry, we expect that there should be even wider classes of models in which scalar, vector, and tensor parts behave differently.   In the present paper, for definiteness we mainly focus on a class of massive gravity models in which propagating degrees of freedom in gravity sector are only two tensor modes with a massive dispersion relation. In this case, the scalar and vector sectors behave exactly like in GR and the only modification emerges in the tensor sector: the dispersion relation of gravitational waves acquires an effective mass, which in general can be time-dependent. We shall discuss an attempt towards a nonlinear completion of such massive gravity models, but in most part of the present paper we shall consider this class of models as a purely phenomenological one to be constrained or probed by observations and experiments. We conduct the analysis based on a general quadratic action for tensor modes, and thus our results apply not only to the theories with precisely two propagating gravitational modes but also to a more general class of theories in which modification of gravity appears primarily in the tensor modes.   Being different from GR only in the tensor sector, observation of gravitational waves and/or their imprints would be the most efficient probe for the model. Therefore, in this paper we address the detectability of the effective mass of gravitational waves by means of the stochastic gravitational wave observations.   The rest of this paper is organized as follows. After describing our model in Sec.~\\ref{Sec:basics}, we discuss an attempt towards nonlinear completion of the model in Sec.~\\ref{Sec:nonlinear}. This consideration at the very least shows that the structure of the model considered in the present paper is not forbidden by the symmetry. However, readers who are interested only in purely phenomenological aspects can safely skip Sec.~\\ref{Sec:nonlinear}. We then derive analytic formulas for the power-spectrum of stochastic gravitational waves and discuss how to read off information about the effective mass of gravitational waves in Sec.~\\ref{Sec:power}.  In Sec.~\\ref{Sec:example} we show some numerical results and make comparison with the analytical results of the previous section. Sec.~\\ref{Sec:Summary} is devoted to a summary of the results and discussions. ",
        "conclusions": "\\label{Sec:Summary}  In this paper, we discussed how to probe the graviton mass using the gravitational wave observations. Our results are based on a general quadratic action given by Eq.~(\\ref{I2}) in which the graviton mass can be time-dependent, and thus they apply universally to any massive gravity theories in which modification from GR appears primarily in the tensor modes. The salient feature of the result is the sharp peak in the gravitational wave spectrum, whose position and height may tell us about the graviton mass of today and the length of the inflationary period. In Sec.~\\ref{Sec:example}, we confirmed our analytical results match well with the numerical results which do not assume the thin-horizon approximation, and also argued the observability of the gravitational wave signal for various observatories. It would be interesting to apply the results to explicit examples of nonlinear completion of massive gravity theories and to derive constraints on such models based on the gravitational wave observations.   Although the example discussed in Sec.~\\ref{Sec:nonlinear} (as well as in Ref.~\\cite{Gumrukcuoglu:2011zh}) was recently shown to suffer from nonlinear instabilities~\\cite{DGM}, it is expected that other working examples exist. A construction with no extra polarization with respect to GR was discussed in \\cite{Dubovsky:2004sg}, where a constant mass contribution appears in the dispersion relations of the gravitational waves. Another potential example is the scenario considered in \\cite{deRham:2012kf}, where the scalar mode is rendered redundant by introduction of some symmetry. Although the vector graviton degrees are still present in this construction, the lack of the scalar graviton would relax the bounds from solar system tests.   Below, we briefly discuss other probes which are sensitive to $M_{GW}(t)$ at the time other than today. One interesting probe is the cosmic microwave background (CMB), especially its B-mode polarization spectrum which is sourced solely by gravitational waves. The signatures in CMB spectrum in the case of the constant mass $M_{GW}$ was discussed by Ref.~\\cite{Dubovsky:2009xk}. Following the argument to derive the B-mode spectrum and generalizing it to the case of time-dependent mass term, we find that the main contribution to the spectrum is sensitive to the value of the mass at the time of the recombination% \\footnote{We find the contribution of the gravitational wave to the CMB B-mode spectrum is roughly proportional to $\\int^{\\eta_r} \\left(k^2+a^2M_{GW}(\\eta)\\right)d\\eta$, where $\\eta_r$ is the conformal time at the time of the recombination. In a generic case, this value is sensitive to $M_{GW}(\\eta_r)$ rather than $M_{GW}(\\eta)$ at any other $\\eta$.}. The most interesting observational feature will be obtained when the graviton mass at that time is in the range $(10\\text{Mpc})^{-1}< M_{GW} < (10\\text{Kpc})^{-1}$. For such a value of $M_{GW}$, the contribution to the B-mode spectrum from the gravitational wave increases compared to the GR case for the modes of low angular multipole $\\ell$, and a plateau in the spectrum up to $\\ell\\lesssim 10^{-3}\\times M_{GW}/H_0 \\sim 100$ is expected to emerge. There will be no signal for smaller values of $M_{GW}$, while for larger values, the entire spectrum will be suppressed, since the rapid oscillation of the gravitational wave, driven by the large $M_{GW}$, will average out the contribution to the CMB polarization spectrum.     Another interesting probe of the modification in tensor sector of gravity is the inspiralling compact binaries; the changes in the propagation speed due to the mass term will give rise to a modification in the evolution of the phase of the gravitational waves emitted by those binaries. Such an effect and the constraint on the constant graviton mass are argued in Ref.~\\cite{Will:1997bb}. It may be interesting to generalize this analysis to the case of time-dependent graviton mass and discuss its observability for the future-planned gravitational wave observatories.  Finally, we give a brief comment on the time evolution of the gravitational wave energy density and a possible constraint on it.% \\footnote{ See Ref.~\\cite{Dubovsky:2004ud} for the arguments on possible production of abundant gravitational wave energy density based on an action similar to our low-energy action of Eq.~(\\ref{I2}). } The energy spectrum of the gravitational wave defined with respect to wavenumber $k$ is given by \\begin{equation} \\Omega_{GW}(k,t) = \\frac{1}{\\rho_\\text{crit}}\\frac{d\\rho_{GW}}{d\\log k} = \\frac{k^3}{12\\pi H^2} \\left( N^{-2}\\dot\\gamma_k^2 + \\omega^2 \\gamma_k^2 \\right) \\propto k^3 \\omega^2 \\gamma_k^2, \\end{equation} where $\\omega$ is given by Eq.~(\\ref{eqn:dispersion-relation}) and the final expression is valid for the WKB solution of Eq.~(\\ref{WKBsol}). The ratio of the energy spectrum of the massive gravity relative to GR can be obtained following the derivations similar to Sec.~\\ref{Sec:power}. To discuss the constraint on the gravitational wave energy density coming from the Big Bang nucleosynthesis (BBN), we need to do the analysis taking $t$ to be the time of BBN. We assume that the gravitational wave primordial amplitude is given by Eq.~(\\ref{primordialAmplitude}), and consider the case that $M_{GW}(t)$ becomes larger than the Hubble scale before BBN. In such a case, the ratio is given by \\begin{equation} \\frac{\\Omega_{GW}(k)}{\\Omega_{GW}^{GR}(k)} = \\begin{cases} 1 & (k_0<k)\\\\ k_0/k & (k_c<k<k_0) \\\\ k_0 / k_c & (k<k_c) \\end{cases} ~, \\end{equation} where $k_0$ and $k_c$ are defined in a similar manner with Sec.~\\ref{Sec:power} replacing $t_0$ by $t_\\text{BBN}$. Roughly speaking, the total energy density of the gravitational wave becomes larger than that in GR by a factor of $k_0/k_c = M_{GW}(t_\\text{BBN})/\\sqrt{H_\\text{BBN}M_{GW}(t_c)}$ if $k_c$ is sufficiently large. Since the total energy density must satisfy $\\int d(\\log k)\\Omega_{GW}(k)<10^{-5}$ at least~\\cite{Maggiore:1999vm}, this enhancement of the gravitational wave energy density may give a tighter constraint on the primordial amplitude of the gravitational wave or the inflationary models from which it is derived. Obviously this constraint will be sensitive to $M_{GW}(t_\\text{BBN})$. It may be interesting to pursue this issue using explicit examples of the massive gravity theories."
    },
    "1208/1208.4648.txt": {
        "abstract": "We present high S/N, moderate resolution ($R\\approx2000$) near-infrared spectra, as well as 10~\\micron\\ imaging, for the brightest members of the central stellar cluster in the W40 \\ion{H}{2} region, obtained using the SpeX and MIRSI instruments at NASA's Infrared Telescope Facility (IRTF). Using these observations combined with archival Spitzer Space Telescope data, we have determined the spectral classifications, extinction, distances, and spectral energy distributions for the brightest members of the cluster.   Of the eight objects observed, we identify four main sequence (MS) OB stars (one late-O, 3 early-B), two Herbig Ae/Be stars, and two low-mass young stellar objects (Class~II).  Strong \\ion{He}{1} absorption at 1.083~\\micron\\ in the MS star spectra strongly suggests that at least some of these sources are in fact close binaries.  Two out of the four MS stars also show significant infrared excesses typical of circumstellar disks.  Extinctions and distances were determined for each MS star by fitting model stellar atmospheres to the SEDs.  We estimate a distance to the cluster of between 455 and 535~pc, which agrees well with earlier (but far less precise) distance estimates.  We conclude that the late-O star we identify is the dominant source of LyC luminosity needed to power the W40 \\ion{H}{2} region and is the likely source of the stellar wind that has blown a large ($\\approx 4$~pc) pinched-waist bubble observed in wide field mid-IR images.  We also suggest that 3.6~cm radio emission observed from some of the sources in the cluster is likely not due to emission from ultra-compact \\ion{H}{2} regions, as suggested in other work, due to size constraints based on our derived distance to the cluster.  Finally, we also present a discussion of the curious source \\irsthreea, which has a very strong mid-IR excess (despite its B3 MS classification) and appears to be embedded in a dusty envelope roughly 2700~AU in size. ",
        "introduction": "\\label{Intro}  \\wforty\\ is a thermal radio continuum source associated with the molecular cloud G~28.74+3.52 in the Aquila Rift~\\citep{Westerhout:1958fk,Zeilik:1978uq}. A faint, patchy optical \\ion{H}{2} region (Sh~64) is  associated with the radio emission~\\citep{Sharpless:1959}.   Radio observations imply a total Lyman continuum luminosity of $\\sim1.5\\times10^{48}$ photons~s$^{-1}$~\\citep{Goss:1970vn,Smith:1985}---about 10\\% of the ionizing flux found in the Trapezium Cluster in Orion---corresponding to  several early B stars or a single O9V.  An IRAS source, cold ammonia core, and a number of mm-wave sources are also associated with the cluster~\\citep{Molinari:1996dz,Maury:2011}, suggesting on-going star formation from dense molecular material.  In addition, \\citet{Crutcher:1982wd} and \\citet{Wu:1996} have found multiple CO line components and/or extended line wings toward \\wforty\\ which suggests outflows driven by young stellar objects (YSOs).   Early infrared maps of the central 3\\arcmin\\ of \\wforty\\ revealed 7 bright sources with optical counterparts behind $9 - 10$ magnitudes of visual extinction~\\citep{Crutcher:1982wd,Smith:1985}.  Spectral energy distributions (SEDs) from IR through millimeter wavelengths suggest substantial circumstellar material around most of these bright sources~\\citep{Smith:1985,Vallee:1994ai}.  2MASS images of the region reveal a cluster of near-IR sources within the central 5\\arcmin\\ of the brightest sources (Fig.~\\ref{fig:W40_2MASS_K_cluster}).  Recent radio observations at 3.6~cm show a cluster of compact sources in the central portion of the W40 IR cluster, many of which correspond to the known infrared sources~\\citep{Rodriguez:2010}.  X-Ray observations of the W40 cluster with the Chandra X-Ray Observatory (CXO) reveal approximately 200 sources associated with the cluster,  the majority of which are thought to be low-mass young stellar objects~\\citep{Kuhn:2010}.  TheX-Ray luminosity function (XLF) obtained by the Chandra  Orion Ultra-deep Project (COUP) indicates a completeness limit of 70\\% (down to 0.1~\\msun) and suggests a total cluster population of $\\approx$600 objects~\\citep{Kuhn:2010}.  UKIRT near-IR observations presented in the same study show that a significant fraction of the X-ray sources also have K-band excesses indicative of circumstellar disks and suggest an age of $< 1$ Myr~\\citep{Kuhn:2010}.  {\\it Herschel} SPIRE and PACS far-IR maps of the entire Aquila Rift region (which includes the Serpens star forming regions and MWC297/Sh-62) reveal $\\approx$201 sources, most of which appear to be associated with W40~\\citep{Bontemps:2010}.  Analysis of the far-IR colors suggest that 45--60 of the detected sources are deeply embedded Class~0 objects~\\citep{Bontemps:2010}.  Depending on the technique used, the distance to W40 ranges from 300 to 900 pc~\\citep[and references therein]{Rodney:2008}, making \\wforty\\ one of  only a few  \\ion{H}{2} regions within a kiloparsec of the Sun.  XLF fitting by \\citet{Kuhn:2010} yields a best-fit distance of $\\approx 600$~pc, while \\citet{Bontemps:2010} adopt a much closer distance of 260~pc based on colocation of W40 with the Serpens region on the sky.  The distance to W40 will be discussed further in Section~\\ref{distance}.  Many questions remain regarding the initial mass function (IMF), age, and star formation history of W40. Even the brightest members of the cluster have not been studied in any detail. In this paper we present infrared observations of these bright members of the W40 IR cluster in an effort to better constrain some of their basic properties, including spectral type, age, and distance.   This is just a first step in better understanding the cluster and star forming region as a whole. ",
        "conclusions": "\\end{deluxetable} %-------------------------------------------   In order to improve our S/N, we re-observed \\irsonea\\ North and South on 13~May and 08--09 July 2010.  Conditions for the May observations were good:  clear skies and $\\sim 0\\farcs7$ seeing in K.  In July conditions were excellent, with photometric skies and seeing from 0$\\farcs$3 -- 0$\\farcs$7 in K.  Observational procedures were the same as for the 2006 observations except with longer integration times and additional pair cycles to improve S/N.  Calibration and reductions were also identical to the 2006 observation run.  Note, however, that the S/N for \\irsoneanorth\\ in the LXD mode was limited by the calibration source which was significantly fainter than \\irsoneanorth .  All data were reduced using Spextool \\citep{Cushing:2004}, the IDL-based package developed for the reduction of SpeX data. The Spextool package performs non-linearity corrections, flat fielding, image pair subtraction, aperture definition, optimal extraction, and wavelength calibration. The sets of spectra resulting from the individual exposures were median combined and then corrected for telluric absorption and flux calibrated using the extracted A0\\, V telluric standard spectra and the technique and software described by \\citet{Vacca:2003}.  Uncertainty in the SpeX flux calibration is $\\sim 10$ \\% \\citep{Cushing:2005,Rayner:2009}.  The spectra from the individual orders were then spliced together by matching the flux levels in the overlapping wavelength regions, and regions of poor atmospheric transmission were removed. The complete, flux calibrated  $0.8-5.0$~\\micron\\ spectra are shown in Figures~\\ref{fig:SEDs_MS} and \\ref{fig:SEDs_YSOs}. %Main Sequence SEDs ----------------------------------- \\begin{figure} \\epsscale{0.49} \\plotone{fig3a.pdf} \\plotone{fig3b.pdf} \\plotone{fig3c.pdf} \\plotone{fig3d.pdf} \\caption{Spectral energy distributions for main sequence sources in W40.  Solid line---SpeX near-IR spectra (this study); Crosses---MIRSI mid-IR photometry (this study); Triangles---Spitzer-IRAC photometry (this study);  Diamonds---\\citet{Smith:1985} photometry; Dotted line---stellar atmospheric model fit for given reddening and distance (see text).} \\label{fig:SEDs_MS} \\end{figure} %---------------------------------------------  %YSO SEDs ----------------------------------- \\begin{figure} \\epsscale{0.49} \\plotone{fig4a.pdf} \\plotone{fig4b.pdf} \\plotone{fig4c.pdf} \\plotone{fig4d.pdf} \\caption{Spectral energy distributions for young stellar objects and pre-main sequence stars in W40.  Line---SpeX near-IR spectra (this study); Crosses---MIRSI photometry (this study); Triangles---Spitzer-IRAC photometry (this study);  Diamonds---\\citet{Smith:1985} photometry.} \\label{fig:SEDs_YSOs} \\end{figure} %---------------------------------------------      \\subsection{MIRSI Observations}  Mid-IR observations of the central portion of the W40 IR cluster were carried out at the IRTF using the Mid-InfraRed Spectrometer and Imager~\\citep[MIRSI,][]{Kassis:2008} instrument  on 8 July 2006 (UT).  MIRSI offers complete spectral coverage over the atmospheric windows at 8-14~\\micron\\ and 18-26~\\micron\\ for both imaging (discrete filters and a circular variable filter) and spectroscopy (in the 10 and 20~\\micron\\ windows with resolutions of 200 and 100, respectively) utilizing a Raytheon 320$\\times$240 pixel Si:As  array with a plate scale of 0$\\farcs$27 per pixel, and a field-of-view of 86\\arcsec$\\times$63\\arcsec\\ at the IRTF. \u00ca   Images of the central 3\\arcmin\\ of the IR cluster were obtained using the 8.7 and 12.3~\\micron\\ filters.  Approximately 15 pointings were used to cover IRS 1 through 3 and IRS 5 with on-source exposure times of 25~s for each pointing in each filter.  Standard chop and nod techniques were used for each pointing to eliminate infrared background emission.  Standard stars $\\alpha$ Lyr (A0V) and $\\mu$ Cep (M2Ia) were observed throughout the night for flux calibration.  The final mosaics are shown in Figure~\\ref{fig:W40_MIRSI}. Note that only the region surrounding IRS~1 -- 3 is shown as IRS 5 was not detected in any of the filters used.  The infrared sources 1E and F, and 2 D -- F are reported here for the first time and have been named according to the convention of \\citet{Smith:1985}.   Fluxes derived using standard aperture photometry techniques are shown in Table~\\ref{table:mirsi_photom}.  For the brighter sources, the uncertainty is dominated by systematic errors in the flux calibration, whereas for the fainter sources, it is dominated by uncertainty in the background emission.  %MIRSI Images ----------------------------------- \\begin{figure} \\epsscale{1.0} \\plottwo{fig5a.pdf}{fig5b.pdf} \\caption{MIRSI mosaic of W40 IRS 1 -- 3 at 8.7 and 12.3~\\micron.  Final mosaic is 2.6\\arcmin $\\times$ 2.3\\arcmin (North is up, East is left).} \\label{fig:W40_MIRSI} \\end{figure} %---------------------------------------------     %For the 24~\\micron\\ flux, the error is dominated by uncertainty in the accepted flux for our standard star $\\mu$~Cep (which is a known variable).   %MIRSI Photometry--------------------------------------- \\begin{deluxetable}{lccl} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{MIRSI photometry} \\tablehead{ \\colhead{Source}  & \\colhead{8.7 \\micron} & \\colhead{12.3 \\micron} & \\colhead{Notes}   \\\\ & \\colhead{mJy (rel. err)} & \\colhead{mJy (rel. err)}& } \\startdata \\irsonea  & 5040.6 (0.06)  & 4104.2 (0.06)    &    \\\\ \\irsoneb  &\t2115.0 (0.1) &\t2602.7 (0.1)  &  \t \\\\ \\irsonec   &\t381.3 (0.1) &\t468.5 (0.1)  & \t \\\\ \\irsoned  &  176.6  (0.1)  &  $<$400     &  \\\\ \\irsonee  &    289.8 (0.1)   &\t 240.8 (0.1)    &  \\\\ \\irsonef  &     364.9 (0.1)  &  1008.8  (0.1)   &  Diffuse \\\\ \\irstwoa    &\t28571.5 (0.06) &\t43307.4 (0.06) &\t17140.2 (0.2) \\\\ \\irstwob    &\t357.1 (0.1) &\t$<$400   & \t\\\\ \\irstwoc    &\t$<$300  &\t\t$<$400   & Not detected\t \t  \\\\ \\irstwod  &  99.1  (0.1)  &   $<$400   \t&  \t \t\\\\ \\irstwoe  &   255.7  (0.1)  &   $<$400   &\t\t \t\\\\ \\irstwof  &  394.4 (0.1)  &   241.1  (0.1)   & Diffuse  \\\\ \\irsthreea  &\t741.9 (0.1) &\t1952.4 (0.1)  &  Resolved, non-point-like  \\\\ \\irsfive       &\t$<$300  &\t\t$<$400   \t& Not detected\t \t  \\\\ \\enddata \\tablecomments{An aperture radius of 3\\arcsec\\ was used for all measurements.} \\label{table:mirsi_photom} \\end{deluxetable} %-------------------------------------------   % Removing for now -- the 1A and B spectra are very low S/N.  2A is good, but not value-added for this paper. %Low-resolution ($R \\approx 200$) spectra were also obtained for \\irsonea , 1b, and 2a in the 10~\\micron\\ window, again using standard chop/nod techniques and total exposure times of 200~s for \\irstwoa\\ and 1b, and 400~s for \\irsonea .  Standard stars $\\alpha$ Lyr (A0V) and $\\mu$ Cep (M2Ia) were observed throughout the night for flux calibration and telluric correction.  The spectrum for each source appears as a smooth continuum with no evidence of silicate emission or absorption or any other narrow features.  \\clearpage   \\label{discussion}  \\subsection{The Distance to W40} \\label{distance}  Depending on the technique used, the distance to the W40 complex has been estimated to be between 300 and 900~pc~\\citep[and references therein]{Rodney:2008,Kuhn:2010}.  If we exclude \\irsthreea\\ (which has a relatively poor fit, and is also exceptional in other ways, see Section~\\ref{sect:irsthreea}), then the best estimate for the distance based on our results (for 3 stars) is between 455~pc and 536~pc.  We can confidently rule out distances less than 340~pc and greater than 686~pc; making it very unlikely that W40 is part of the same complex that includes the Serpens star forming regions, as assumed by~\\citet{Bontemps:2010}.  Our best estimates are somewhat lower than the distance estimate of 600~pc derived from X-ray luminosity function (XLF) fitting method used by \\citet{Kuhn:2010}. (In fact, a distance of 500~pc produces a qualitatively worse fit to the XLF according to the authors.) Our distance determination has the advantage of being based on (admittedly only a few) well-determined spectral types, but only for a few sources that happen to be OB stars with a wide range of possible luminosities.  The advantage of the XLF method is that it is based on statistics over a large number of observed sources;  the disadvantage is that one must assume a ``universal XLF'' that may not apply to all star forming regions.  For the remainder of the paper, we will adopt a distance of 500~pc.  \\subsection{IRS 3A} \\label{sect:irsthreea} As indicated in Section~\\ref{sect:SEDs}, \\irsthreea\\ has significant mid-IR excess with a generally flat spectral index ($\\alpha_{mir} =  -0.3$).   In addition,  \\irsthreea\\ is also {\\em resolved} as a circular source at longer mid-IR wavelengths.  This can be seen readily in both the MIRSI 8.7 and 12.3~\\micron\\ images (FIg.~\\ref{fig:W40_MIRSI}), and in the IRAC channel 4 images (see red channel in Fig.~\\ref{fig:W40_IRAC_134}).  In fact, \\irsthreea\\ is point-like only in the IRAC channel 1 (3.6~\\micron) data;  in the other bands it has a roughly Gaussian profile with a FWHM that exceeds that of the   corresponding PSF.  The measured FWHM as a function of wavelength (with comparison to the PSF FWHM) is shown in Table~\\ref{table:irsthreea_size}. At 500~pc, the measured FWHM at 4.5~\\micron\\ corresponds to a size of 1050~AU, increasing to 2700~AU at 12.3~\\micron.  The fact that \\irsthreea\\ is resolved in the near- to mid-IR indicates that the emission must be coming from a circularly symmetric distribution of dust---either a spherical envelope or a disk viewed roughly face-on--- with a radius of roughly 500~AU.  There are, however, issues with both of these interpretations.  In general, near- to mid-IR emission from circumstellar disks around intermediate mass stars (e.g. HAeBe stars) is generated at the disk surface where dust temperatures exceed 250~K.  For most flaring disk models, this corresponds to regions at the disk surface within 10~AU of the central star~\\citep[e.g.][]{Dullemond:2002}, and hence would be {\\em unresolved} in our observations.  On the other hand, an optically thick spherical envelope would produce additional extinction in the optical that would produce a larger $A_V$ for \\irsthreea\\ than the other stars in W40, which is not observed.  A more probable explanation is that the resolved emission from 4 -- 12 \\micron\\ is due to emission from small grains (transiently heated by the strong UV emission from \\irsthreea) in an optically thin, dusty nebula (or shell) immediately surrounding the star~\\citep[e.g.][]{Hartmann:1993}.  If this is the case, then PAH and [\\ion{Ne}{2}] emission from this dusty nebula could also explain the excess flux observed in the 8.0 and 12.3~\\micron\\ bands (see Figure~\\ref{fig:SEDs_MS}).  This dusty shell may be a relic from an earlier pre-main sequence (Herbig~Be) phase which has not yet been dispersed by the stellar winds of \\irsthreea .   %ORIGINAL:  Since circumstellar disks around Main Sequence stars are typically $< 1000$~AU in size, we posit that the diffuse emission comes from a dusty envelope.  Though rare for Main Sequence B stars, such envelopes have been invoked to explain the large IR luminosity for Herbig~Be objects for some time~\\citep{Waters:1998,Millan-Gabet:2001}.  PAH and [\\ion{Ne}{2}] emission from this dusty envelope, driven by UV radiation from \\irsthreea\\ (a B3 star), could also explain the excess flux observed in the 8.0 and 12.3~\\micron\\ bands (see Figure~\\ref{fig:SEDs_MS}).  This dusty envelope may be a relic from an earlier pre-main sequence (Herbig~Be) phase which has not yet been dispersed by the stellar winds and ionizing radiation of \\irsthreea .   %IRS 3A Size vs. wavelength--------------------------------------- \\begin{deluxetable}{lccc} \\tabletypesize{\\footnotesize} \\tablewidth{0pt} \\tablecaption{IRS~3A:  Size vs. Wavelength} \\tablehead{ \\colhead{Waveband}  & \\colhead{PSF\\tablenotemark{a} }  & \\colhead{IRS 3A}  & \\colhead{IRS 3A}  \\\\  \\colhead{(\\micron)}  & \\colhead{FWHM (\\arcsec)}  & \\colhead{FWHM (\\arcsec)}  & \\colhead{FWHM (AU)\\tablenotemark{b}} } \\startdata IRAC1, 3.6 &\t1.66 &\t1.764 &\t882 \\\\ IRAC2, 4.5 &\t1.72 &\t2.1 &\t1050  \\\\ IRAC3, 5.8 &\t1.88 &\t2.58 &\t1290 \\\\ IRAC4, 8.0 &\t1.98 & 4.158 &\t2079 \\\\ MIRSI, 8.7\t & 1.21 &\t4.32 &\t2160 \\\\ MIRSI, 12.3 &\t1.32  &\t5.4 &\t2700 \\\\ \\enddata \\tablenotetext{a}{IRAC PSF taken from mean FWHM of the point response function (PRF) reported in the IRAC Instrument Handbook.  MIRSI PSF based on measured performance and is nearly diffraction limited.} \\tablenotetext{b}{Assuming $d = 500$~pc.} \\label{table:irsthreea_size} \\end{deluxetable} %-------------------------------------------     The fact that the Kurucz model fits yield a distance to W40 that is far closer than that of the other MS stars in our (admittedly small) sample, and the presence of strong \\ion{He}{1} absorption strongly suggests that \\irsthreea\\ is in fact a close binary.  Increasing the distance of \\irsthreea\\ by a factor $\\approx$1.3 would place the binary at the same distance as \\irstwob\\ and \\irsfive ,  and would result in a  $\\approx$1.7 fold increase in the luminosity, which could be accounted for by a companion of mid-B spectral type. It is also possible that \\irsthreea\\ is overluminous because it is a ``bloated'' star, i.e. the stellar radius is much larger than a typical MS B3~\\citep[e.g][]{Ochsendorf:2011}.  Pre-main sequence evolutionary tracks for intermediate mass stars show a significant increase in total luminosity just before the star settles onto the main sequence~\\citep{Palla:1993,Hosokawa:2009}.  It is possible that \\irsthreea\\ is nearing (or at) this peak in luminosity, which would contribute to the poor stellar model fit, result in a shortened distance, and explain the existence of the extended dusty envelope.  In this case, the strong \\ion{He}{1} absorption could still be due to a much lower mass binary companion.  \\subsection{Compact Radio Emission}  All of the sources we identify as YSOs in our sample are also detected in the radio survey of~\\citet{Rodriguez:2010} (see Table~\\ref{table:spectral_features}).  Only two of the main sequence stars in our study are also included by the radio maps of~\\citet{Rodriguez:2010} (\\irsoneasouth, and \\irstwob) and neither is detected at 3.6~cm.  \\citet{Rodriguez:2010} suggest that the non-variable compact radio sources are likely due to ultra-compact \\ion{H}{2} (UC\\ion{H}{2}) regions around young, high mass stars.  This seems unlikely, however, as the size of the unresolved radio sources at a distance of 500~pc would be less than $\\approx 100$~AU---much smaller than a typical UC\\ion{H}{2} region\\citep{Kurtz:2005}.  It seems more likely that the 3.6~cm continuum flux would be due to free-free emission from shocked gas within 100~AU of the YSO caused by a jet or outflow.  It could also be driven by synchrotron emission associated with the stellar magnetic field, a conclusion that is  supported by the fact that all the YSOs in our sample are also strong X-ray sources (see Table~\\ref{table:spectral_features}).  \\subsection{The W40 \\ion{H}{2} Region}  The predicted Lyman continuum (LyC) fluxes for each of our main sequence sources (based on the \\citet{Kurucz:1993} model) is given in Table~\\ref{table:ms_stars}.  The total LyC flux in the region is dominated by  the O9.5 star \\irsoneasouth\\ at 6.3$\\times 10^{47}$~photons~s$^{-1}$.  Based on our spectral type analysis (Sect~\\ref{analysis:ms_stars}), it is unlikely that \\irsoneasouth\\ is earlier than O9 which places an upper limit on the LyC flux available to drive the \\ion{H}{2} region of $\\approx 1.6 \\times 10^{48}$ photons~s$^{-1}$.  This upper limit agrees very nicely with the derived LyC flux of $\\sim1.5\\times10^{48}$ photons~s$^{-1}$ from radio observations~\\citep{Goss:1970vn,Smith:1985}.  Based on the IR luminosity, \\citet{Smith:1985} originally suggested that \\irstwoa\\ must be the dominant energy source in W40.  We conclude, however, that  \\irsoneasouth\\ must be the main source of ionizing radiation.  The morphoplogy of the \\ion{H}{2} region in the thermal IR is quite different from that seen in the POSS plates and 2MASS images.  Images of the region from the Mid-course Space Experiment (MSX) obtained from the NASA \\anchor{http://irsa.ipac.caltech.edu/Missions/msx.html}{Infrared Science Archive (IRSA)}\\footnote{ \\url{http://irsa.ipac.caltech.edu/Missions/msx.html}} (see Fig.~\\ref{fig:W40_MSX}) reveal an hour-glass shaped structure centered roughly on \\irstwoa, approximately 28\\arcmin\\ in height and about 17\\arcmin\\ in width~\\citep{Rodney:2008}, corresponding to 4.1 by 2.5~pc at a distance of 500~pc.   The long axis of the hourglass is at a PA of $\\sim 300$~degrees (NNW).  The structure is most prominent in the MSX A-band (8.3 ~\\micron) images, where the brightness is likely dominated by emission in the PAH line at 7.7~\\micron. The hourglass does not appear related to the weak CO outflow centered on the dense material roughly 3\\arcmin\\ to the West~\\citep{Zhu:2006fu}.  We interpret this double-lobed morphology as a pinched waist ``bubble''  blown by a symmetric wind, most likely from the O9.5 star \\irsoneasouth.  The waist axis of the hourglass lies along a PA of roughly 60~degrees (ENE).  The Spitzer IRAC images (Fig.~\\ref{fig:W40_IRAC_134}) reveal numerous bright-rimmed sheets and clumps (``elephant trunks'') along this axis, nearly all of which point back toward \\irsonea, further underscoring \\irsoneasouth\\ as the dominant wind and UV source of the region.   Far-IR and mm-wave observations of W40 reveal 36 deeply embedded sources (prestellar cores, and Class 0,I objects) indicating that star formation is on-going, and has perhaps been initiated by the pressure wave associated with the wind-bubble~\\citep{Maury:2011}.     %W40 MSX Image ----------------------------------- \\begin{figure} \\epsscale{0.5} \\plotone{fig10.pdf} \\caption{MSX 3-color composite of W40 at 8.3 (blue), 12.1 (green), and 14.6~\\micron\\ (red).  Field of view is approximately 30\\arcmin.  Red circle marks the center of the peak CO emission \\citep{Zhu:2006fu}. } \\label{fig:W40_MSX} \\end{figure} %---------------------------------------------    %In the 2MASS K-band images there is a  circular patch of nebulosity surrounding \\irsfive\\ roughly 70\\arcsec\\ across, corresponding to a diameter of $\\sim 0.2$~pc at a distance of 500~pc.  Furthermore, there appears to be a cavity surrounding \\irsfive\\ in the Spitzer IRAC images.  These features suggest that \\irsfive\\ was formed in situ, and perhaps separately from the cluster of stars surrounding IRS 1 and 2.  \\irsfive\\ has the least amount of mid-IR excess of all the main sequence stars in our sample, suggesting that it is the oldest of the group.  \\clearpage"
    },
    "1208/1208.6001_arXiv.txt": {
        "abstract": "Recent observations and simulations have suggested that \\HII regions around massive stars may vary in their size and emitted flux on timescales short enough to be observed. This variability can have a number of causes, ranging from environmental causes to variability of the ionizing source itself. We explore the latter possibility by considering the pre-main-sequence evolution of massive protostars and conducting numerical simulations with ionizing radiation feedback using the \\FLASH AMR hydrodynamics code. We investigate three different models: a simple ZAMS model, a self-consistent one-zone model by Offner et al. (2009), and a model fit to the tracks computed by Hosokawa \\& Omukai (2009). The protostellar models show that hypercompact \\HII regions around massive, isolated protostars collapse or shrink from diameters of 80 or 300 AU, depending on the model choice, down to near absence during the swelling of stellar radius that accompanies the protostar's transition from a convective to a radiative internal structure. This occurs on timescales as short as $\\sim$3000 years. ",
        "introduction": "Massive stars dominate their birth environments through a set of powerful feedback processes, the full implications of which are not completely understood. Among their various feedback mechanisms is the formation of \\HII regions---bubbles of hot ($10^4$ K), ionized gas that expand into the colder ($10^2$ K) surrounding medium. These regions are fueled by the prodigious amounts of UV radiation emitted by massive stars \\citep{Hoare+2007,Beuther+2007}. By heating and ionizing the gas around them, massive stars alter their birth environments and provide detectable observational signatures.  \\HII regions may contribute to the destruction of molecular clouds \\citep{Keto2002,Keto2003,Keto2007,Matzner2002}, helping to set cloud lifetimes. They are observable by their radio continuum emission \\citep{MezgerHenderson1967}, or by their recombination lines (e.g.\\ \\citet{WoodChurchwell1989} use the H76$\\alpha$ line). Even Young Stellar Objects (YSOs) emit enough UV radiation to begin ionizing the gas around them \\citep{Bik+2005,Churchwell2002}, which might teach us something about the lives of massive protostars before they reach the main sequence.  More recently, observations have shown time variability in \\HII regions \\citep{Franco-HernandezRodriguez2004,Rodriguez+2007,Galvan-Madrid+2008,gomezetal08}. \\citet{Franco-HernandezRodriguez2004} have suggested that such observed time-variability may be due to the changes occurring in the source of the ionizing radiation, though it may also be due to increased absorption in the rapidly-evolving core of the nebula.  Three-dimensional collapse simulations of massive star formation that included feedback by ionizing radiation showed time variability with changes in size and flux very similar to the observations \\citep{Peters2010a,Klassen+2012}. The origin of this time variability is the strong accretion flow around the massive star, which is dense enough to shield the ionizing radiation. As a consequence, the \\HII region fluctuates between extended and trapped states as long as the star is embedded in a sufficiently strong accretion flow. In the simulations, the flickering stops when companion stars form in the gravitationally unstable accretion flow around the central massive star and intercept material that would otherwise be accreted onto the high-mass star, a process we call ``fragmentation-induced starvation'' \\citep{Peters2010c}. Then the accretion flow becomes weak enough that the \\HII region can isolate the forming high-mass star and terminate the accretion process entirely. Magnetic fields, which provide additional support against gravitational collapse and reduce the number of companion stars formed, do not seem to directly affect the \\HII region variability \\citep{Peters2011}.  The time variability during the accretion phase can lead to morphological changes of the \\HII region appearance over timescales as short as $\\sim 10$ years and appears able to resolve the lifetime problem for ultracompact \\HII regions \\citep{WoodChurchwell1989,Peters2010b}. With the aid of a statistical analysis of the synthetic radio continuum observations of the \\HII regions formed in the simulations, \\citet{Galvan-Madrid+2011} predicted that 7\\% of ultracompact and hypercompact \\HII regions should have detectable flux increments, and 3\\% should have detectable flux decrements when observed at two epochs separated by about 10 years. Particularly interesting are shrinking \\HII regions, for which \\citet{Galvan-Madrid+2011} found that flux decrements by at least 10\\% are twice as likely as flux decrements by at least 50\\% for this time interval.  In general, however, \\HII region variability has two possible causes. The first is the chaotic gas motions around a forming star, where denser, optically thick gas can intercept stellar radiation and causing the shadowed gas to neutralize \\citep{Peters2010a}. The second possible source of \\HII region variability is due to the star itself, as already noted by \\citet{Franco-HernandezRodriguez2004}. Pre-main-sequence stars undergo changes during their early evolution, sometimes swelling and other times contracting \\citep{YorkeBodenheimer2005}. This can potentially occur on even short time scales ($<2000$ years). Changes in the surface temperature of a star will affect its UV flux, which in turn affects the size of the \\HII region.  The connection between protostellar variability and \\HII region variability has not yet been explored. To do so, simulations must be equipped with good protostellar models. We ask: can early variability of an \\HII region be traced directly to the pre-main-sequence evolution of the massive protostar itself? Models of prestellar evolution have been investigated by \\citet{PallaStahler1991}, \\citet{PallaStahler1992}, \\citet{Nakano2000}, \\citet{McKeeTan2003}, \\citet{Offner2009} and \\citet{HosokawaOmukai2009}, among others, but never investigating the \\HII region variability that could result from a massive evolving protostar.  In a previous paper \\citep{Klassen+2012}, we simulated the effects of using prestellar models in simulations of clustered massive star formation with initial conditions as in \\citet{Peters2010a}. Fluctuations in \\HII region size and morphology due to the environmental factors were unaffected by the choice of prestellar model, as might be expected. The prestellar model we implemented, however, did show a delayed onset of major heating and ionization in the cluster.%  In the present paper, we investigate the effect of pre-main-sequence evolution on the ionization feedback of a massive protostar with radiation-hydrodynamic simulations that follow the formation and expansion of an \\HII region around a single, isolated star embedded in a homogeneous environment. Our reasoning for resorting to a simple model calculation instead of running a full collapse simulation as in \\citet{Klassen+2012} is twofold. First, it allows us to separate time variability of the \\HII region due to absorption of ionizing radiation by the accretion flow from time variability caused by structural changes in the protostar. In a realistic simulation, both effects would occur simultaneously, making it difficult to assess their relative importance. Second, the prestellar model by \\citet{Offner2009}, which was used in our collapse simulations \\citep{Klassen+2012}, represents the swelling phase of the more detailed calculations by \\citet{HosokawaOmukai2009} only very crudely. Following the behavior in \\citet{HosokawaOmukai2009} more closely allows for a better modeling of the time variability. However, accretion rates in collapse simulations are strongly fluctuating \\citep{Peters2010a,Klassen+2012}, as opposed to the constant accretion rates considered by \\citet{HosokawaOmukai2009}.  Thus, in order to create a realistic model for the time variability caused by pre-main-sequence evolution of the high-mass protostar, we are forced to use idealized simulations with a constant accretion rate as the ones presented in this paper. The work is further idealized because we deliberately only explore constant accretion rates, and thus neglect any time dependence in the HII region due to changing accretion. As a result of this simplicity, we are able to present a clear comparison of three sub-grid models for protostellar effects.  It is at present not possible to run a detailed pre-main-sequence evolution code along with a hydrodynamical simulation, so that our approach of investigating the two aspects of the problem separately appears to be the only solution.  We also stress that we cannot, in general, obtain the time dependence of the \\HII region radius by analytical arguments. The key point is that the radius of the protostar may grow or contract depending on the stellar evolutionary stage, which in turn also depends on the accretion history. Since the ionizing flux grows initially during the \\HII region expansion, both an enhanced ionizing flux and the thermal pressure of the already ionized gas work together to extend the \\HII region radius, so that the \\HII region can neither be modeled by a Str\\\"{o}mgren sphere nor by a standard D-type ionization front. Similar considerations hold for the shrinking phase, in which the ionized gas partially recombines.  Generally, a changing stellar radius affects the surface temperature, which in turn controls the UV flux. This means that if the effective temperature can change significantly over a short period of time, changes in the observed flux or size of \\HII regions may signal evolutionary changes taking place in the stars themselves, whose ionizing radiation is driving the region.  Our numerical approach is described in Section \\ref{sec:numerical_methods}. In Section \\ref{sec:results} we show our results for the evolution of a single massive protostar in a uniform medium, accreting matter at a constant rate, comparing the effects of different accretion rates and different prestellar models. Our assessment of these effects is discussed in Section \\ref{sec:discussion}, with some considerations of the observability of \\HII region variability. Our findings are summarized in \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Our results show that massive protostars in uniform media emit ionizing radiation capable of forming a hypercompact \\HII region. In our simulations using the \\FLASH astrophysics code with ionizing feedback and different prestellar models, the pre-main-sequence evolution of protostars can cause early variability in the size of these \\HII regions. During pre-main-sequence evolution, the stellar radius swells dramatically when the star transitions to a radiative internal structure. This is accompanied by a cooling of the surface temperature and a drop in the ionizing flux by several orders of magnitude. During this transition, young \\HII regions will shrink or collapse until the contracting star has regained its previous ionizing flux.  Our results show that pre-main-sequence evolution can affect the growth of \\HII regions around massive protostars on potentially short time scales. The collapse of the \\HII region, initially $\\sim$300 AU across in the \\citet{Offner2009} model, happened within a single simulation timestep (8 years) when the radius of the star was enlarged during the evolution. The \\citet{HosokawaOmukai2009} model saw the radius of the \\HII region contract from an initial diameter of $\\sim$90 AU to zero over the course of about 3.5 kyr.  The observability of this variability may be possible with the best interferometric radio observations such as with VLBI, the EVLA, or ALMA. Nevertheless, model simulations of hypercompact \\HII regions are an important part of understanding the complete picture of massive star formation. Massive protostars are unique in that their UV fluxes are large enough to ionize gas even before the star has reached the main sequence, and it is therefore important that these early states be further explored."
    },
    "1208/1208.1767_arXiv.txt": {
        "abstract": "There are only a few tracers available to probe the kinematics of individual early-type galaxies beyond one effective radius. Here we directly compare a sample of planetary nebulae (PNe), globular clusters (GCs) and galaxy starlight velocities out to $\\sim$4 effective radii, in the S0 galaxy NGC 2768. Using a bulge-to-disk decomposition of a K-band image we assign PNe and starlight to either the disk or the bulge. We show that the bulge PNe and bulge starlight follow the same radial density distribution as the red subpopulation of GCs, whereas the disk PNe and disk starlight are distinct components. We find good kinematic agreement between the three tracers to several effective radii (and with stellar data in the inner regions). Further support for the distinct nature of the two galaxy components come from our kinematic analysis. After separating the tracers into bulge and disk components we find the bulge to be a slowly rotating pressure-supported system, whereas the disk reveals a rapidly rising rotation curve with a declining velocity dispersion profile. The resulting V$_{rot}$/$\\sigma$ ratio for the disk resembles that of a spiral galaxy and hints at an origin for NGC 2768 as a transformed late-type galaxy. A two-component kinematic analysis for a sample of S0s will help to elucidate the nature of this class of galaxy. ",
        "introduction": "Based on studies of the morphology-density relation in clusters (Dressler 1980) and groups (Wilman et al. 2009), it has been suggested that lenticular (S0) galaxies may be the descendants of spirals that have undergone some evolutionary process (e.g. ram pressure stripping, galaxy harrassment, gas starvation and/or mergers). Recent investigations of S0s have studied their metallicity gradients (Bedregal et al. 2011), Tully-Fisher relation (Bedregal, Aragon-Salamanca \\& Merrifield 2006) and stellar populations (Aragon-Salamanca, Bedregal \\& Merrifield  2006). However, the origin of S0s is still a subject of much debate (e.g. Kormendy \\& Bender 2012). The internal kinematics of galaxies are a key tool to understanding their structure and formation histories, and S0s are no exception. For example, the kinematics will be largely unaffected if S0s were formerly spirals that have been simply stripped of gas or if they were involved in a relatively minor merger.   Although half of the stellar mass within a galaxy lies within one effective radius (R$_e$), more than 90\\% of the {\\it total} mass and angular momentum does not (Romanowsky \\& Fall 2012). Thus in order to examine the internal kinematics and total mass of early-type galaxies one must probe well beyond 1 R$_e$. elliptical and S0 (early-type) galaxies often lack the significant quantities of extended HI gas commonly found in spirals, so the kinematic tracers are the galaxy starlight itself, planetary nebulae (PNe) and globular clusters (GCs).    Using the underlying starlight of a galaxy to probe its stellar kinematics is perhaps the preferred method, however the surface brightness of a galaxy declines rapidly with increasing radius so it is very difficult to obtain high quality spectra beyond a few effective radii without a large investment of 8m-class telescope time (Coccato et al. 2010) or using deep single pointings (Weijmans et al. 2009). PNe and GCs have the advantage that they are ubiquitous in the halos of early-type galaxies out to large galactocentric radii (5--10 R$_e$). Although there have been several studies of PNe and GC system kinematics in early-type galaxies, very few studies have directly compared them to each other, or to results from galaxy starlight over a common radial range.  Luminous PNe are the end product of low mass stars. However, there is still debate as to whether they arise from normal single-star evolutionary processes or from mass transfer in a binary star system (Ciardullo et al. 2005) In the former case, the PNe observed in early-type galaxies would have an age of $\\sim$1.5 Gyr, and in the latter case they could be as old as 10 Gyr. Coccato et al. (2009) showed that the radial surface density of PNe follows the galaxy starlight in early-type galaxies and that they are useful probes of galaxy kinematics. However, ellipticals with embedded thick disks and S0 galaxies may contain two subpopulations of PNe, one associated with the disk and one with the bulge, as seen for spiral galaxies (Nolthenius \\& Ford 1987; Hurley-Keller et al. 2004). Not accounting for the different kinematics of these distinct PNe subpopulations could lead to misleading results as illustrated recently by Cortesi et al. (2011) for the lenticular galaxy NGC 1023. Furthermore, Dekel et al. (2005) suggested that an intermediate-aged population of PNe may have `contaminated' the PNe velocity dispersions of some galaxies in the early-type sample of Romanowsky et al. (2003) and hence impacted the resulting mass analysis.  The globular cluster (GC) systems of all large galaxies, irrespective of Hubble type, generally consist of two subpopulations -- blue (or metal-poor) and red (or metal-rich). Both of these subpopulations are thought to have ages $\\ge$ 10 Gyr and hence trace old stellar populations (for a review of GC system properties see Brodie \\& Strader 2006). The blue subpopulation is associated with galaxy halos (Forte et al. 2005; Forbes et al. 2012) whereas the red subpopulation has been shown to share many properties with the spheroid/bulge stars of early-type galaxies (Strader et al. 2011; Forbes et al. 2012), including their kinematics (Pota et al. 2012). We note that the association of red GCs with the bulge and not with a thin disk component extends to spiral galaxies, including our own  (Minniti 1995; Cote 1999) and the Sombrero (Forbes, Brodie \\& Larsen 2001).    To better understand the issues discussed above it is important to directly compare different kinematic tracers for the same galaxy. Here we combine starlight, PNe and GC data for an archetype lenticular galaxy NGC 2768 and directly compare these different kinematic tracers in the same galaxy. The galaxy is a nearby, near edge-on S0 (Sandage, Tammann \\& van den Bergh 1981), although we note that it was originally classified as an E6 in the RC3 catalogue. According to Wikland et al. (1995) is it an isolated galaxy however it has also been classified as part of the Lyon Group of Galaxies (LGG) 167 (Garcia 1993). It reveals ionised gas and a minor axis dust lane (Kim 1989). The central ionised gas and stars are known to have different kinematics (Fried \\& Illingworth 1994) suggesting an external origin for the gas. NGC 2768 is a rare example of an early-type galaxy with detectable CO emission (Wikland et al. 1995) and the host of a Calcium-rich supernova type Ib (Filippenko \\& Chornock 2000). The effective radius of the galaxy is 1.06 arcmin (Proctor et al. 2009; Cappellari et al. 2011). For a distance of 21.8 Mpc (Cappellari et al. 2011), this corresponds to 6.7 kpc. ",
        "conclusions": "By combining galaxy starlight, PNe, and red GC data for NGC 2768 we have shown that the red subpopulation of GCs and a subpopulation of PNe follow the same radial surface density profile as the bulge component of the galaxy starlight. An additional distinct, and presumably younger, PNe subpopulation is associated with the galaxy disk.  We have also investigated the radially extended kinematics of these three tracers out to $\\sim$ 4R$_e$ and show that they are in very good agreement with each other over a large range of radii. Using a 2D probability map generated from a K-band bulge-to-disk decomposition we estimate the bulge-to-total ratio to be 0.7, with both components having a similar scale size. Using this two-component decomposition we examined the kinematics of the bulge and disk separately finding that the disk component reveals rapid rotation and a falling velocity dispersion profile resulting in an increasing V$_{rot}$/$\\sigma$ ratio with radius. Such a behaviour is similar to that seen in the disks of late-type spiral galaxies. We also found a slowly rotating bulge, indicative of a pressure-supported system.  Although NGC 2768 is an S0 galaxy, many elliptical galaxies contain embedded disks and potentially an associated subpopulation of disk PNe. This work shows that bulge and disk components may need to be taken into account in any kinematic, or mass, analysis of early-type galaxies. This is particularly true when using kinematics to help unravel the origin of S0 galaxies."
    },
    "1208/1208.3281_arXiv.txt": {
        "abstract": "Using a mid-infrared calibration of the Cepheid distance scale based on recent observations  at 3.6~$\\mu$m with the {\\it Spitzer Space Telescope}, we have obtained  a new, high-accuracy calibration of the Hubble constant. We have established the mid-IR zero point of the Leavitt Law (the Cepheid Period-Luminosity relation) using time-averaged 3.6~$\\mu$m data for ten high-metallicity, Milky Way Cepheids having independently-measured trigonometric parallaxes.  We have adopted the slope of the PL relation using time-averaged 3.6~$\\mu$m data for 80 long-period Large Magellanic Cloud (LMC) Cepheids falling in the period range 0.8 $<$ log(P) $<$ 1.8.  We find a new reddening-corrected distance to the LMC of 18.477 $\\pm$ 0.033 (systematic) mag. We re-examine the systematic uncertainties in $H_0$, also taking into account new data over the past decade. In combination with the new {\\it Spitzer} calibration, the systematic uncertainty in $H_0$ over that obtained by the {\\it Hubble Space Telescope} ({\\it HST}) {\\it Key Project} has decreased by over a factor of three. Applying the {\\it Spitzer} calibration to the {\\it Key Project} sample, we find a value of $H_0$ = 74.3 with a systematic uncertainty of $\\pm$ 2.1 (systematic)~km s$^{-1}$ Mpc$^{-1}$, corresponding to a 2.8\\% systematic uncertainty in the Hubble constant. This result, in combination with WMAP7 measurements of the cosmic microwave background anisotropies and assuming a flat universe, yields a value of the equation of state for dark energy, $w_0$ = -1.09 $\\pm$ 0.10. Alternatively, relaxing the constraints on flatness and the numbers of relativistic species, and combining our results with those of WMAP7, Type Ia supernovae and baryon acoustic oscillations yields $w_0$ = -1.08 $\\pm$ 0.10 and a value of N$_{eff}$ = 4.13 $\\pm$ 0.67, mildly consistent with the existence of a fourth neutrino species. ",
        "introduction": "Over the past several decades, there has been a steady increase in the accuracy with which extragalactic distances and the Hubble constant can be measured ({\\it e.g.,} Freedman {\\it et al}. 2001, hereafter F01; Riess {\\it et al}. 2011; Komatsu {\\it et al}. 2011; and for a review see Freedman \\& Madore 2010). This has resulted from a number of factors: the availability, on the ground and especially in space, of high throughput, high dynamic range optical CCDs and infrared arrays; the multi-wavelength sensitivity of these devices; making it possible to correct for systematic effects of reddening and metallicity; and to use a wider range of methods for measuring relative distances beyond the immediate reach of Cepheid variables. Just over a decade ago, there was still debate over the value of the Hubble constant at a level of a factor of two; today, there is the promise of measuring the Hubble constant to an accuracy better than two percent.    In combination with other constraints ({\\it e.g.,} the angular power spectrum of cosmic microwave background anisotropies), an independent measurement of H$_0$ to accuracy of better than a few percent can provide critical constraints on the dark energy equation of state, the spatial curvature of the universe, neutrino physics and general relativity (see Suyu {\\it et al}. 2012 for a recent discussion).  In practice, an accurate value of H$_0$ provides a means of breaking the degeneracies amongst several cosmological parameters. For example, measurements of cosmic microwave background anisotropies yield well-determined values of the products of $\\Omega_m h^2$, $\\Omega_b h^2$ (where $\\Omega_m$ and $\\Omega_b$ are the matter and baryon densities, respectively), but not the densities, nor H$_0$ independently.  There are other degeneracies between H$_0$ and the equation of state, w$_0$, as well as its evolution, w$_a$; between H$_0$ and the number of relativisitic neutrinos, N$_{eff}$ and the sum of the masses of neutrinos; and between H$_0$ and $\\sigma_8$, the fluctuation of matter on 8 Mpc scales (for recent discussions, see Dunkley {\\it et al}. 2009; Komatsu {\\it et al}. 2009). Hence, the motivation for measuring an independent value of H$_0$ accurate to a few percent has continued to increase.   As described in Freedman {\\it et al}. (2011, hereafter F11), we have begun a new Carnegie Hubble Program (CHP), specifically designed to minimize and/or eliminate the remaining known systematics in the measurement of the Hubble constant using mid-infrared data from NASA's {\\it Spitzer Space Telescope}. Here we report on a newly derived value of the Hubble constant and its uncertainty, based on the first data acquired from this program for Cepheids in the Milky Way and the LMC. Our focus in this paper is primarily the zero point of the Cepheid extragalactic distance scale, and a reassessment of the systematic error budget. ",
        "conclusions": "As we saw in Figure 3, we have graphically summarized  the decrease in each of the systematic uncertainties, comparing the {\\it HST} {\\it Key Project} and the CHP.  There are four key systematic improvements to the Cepheid distance scale that have occurred in the decade separating this study and the {\\it HST} {\\it Key Project.}  First, the {\\it Spitzer} 3.6~$\\mu$m data provide a zero point that is about an order of magnitude less sensitive to total line-of-sight extinction than the optical bands used for the {\\it Key Project}. Systematic uncertainties in the reddening corrections (and uncertainties in the extinction law itself) are virtually eliminated in moving to the mid-IR.  Second, longer-wavelength data are theoretically predicted ({\\it e.g., } Marconi {\\it et al}. 2005; Romaniello {\\it et al}. 2008) and empirically demonstrated (Freedman \\& Madore 2010; Freedman {\\it et al}. 2011; Riess {\\it et al}. 2011; this paper) to be less sensitive to metallicity. Moreover, the Milky Way Cepheid sample (now setting the zero point) has a metallicity more comparable to those of the majority of the {\\it HST} {\\it Key Project} spiral galaxies, thereby eliminating the bulk of the systematic uncertainty involved in previously using the (lower-metallicity) LMC Cepheids for the zero-point calibration.  Third, there are now direct parallax measurements for a representative sample of Milky Way Cepheids to define geometrically the absolute zero point of the Leavitt relation (Benedict {\\it et al}.  2007).  This new {\\it Spitzer} Galactic trigonometric zero point eliminates the long-standing dependence on the distance to the LMC. Long-period LMC Cepheids simply define the slope and width of the 3.6~$\\mu$m Leavitt law. In addition, independent geometric methods for measuring the distance to the LMC agree with the Cepheid calibration to within 1.5\\% $rms$ in distance, providing an external check on this new calibration.  Fourth, and independently of the CHP, new near-infrared Cepheid distances to galaxies containing Type Ia supernovae (Riess {\\it et al}. 2011), as well as larger samples of more distant supernovae (Hicken {\\it et al}. 2009) have become available.  Based on our analysis of the {\\it Spitzer} data available to date, combined with data from the Hubble {\\it Key Project}, we find $H_0$ = 74.3 $\\pm$ 0.4 (statistical) $\\pm$ 2.1 (systematic) km s$^{-1}$ Mpc$^{-1}$. This value of $H_0$ is in excellent agreement with that of the {\\it Key Project}, as well as that of Riess {\\it et al}. (2011). Combining this result with constraints from WMAP7 alone yields a value for the dark energy equation of state of $w_0$ = -1.09 $\\pm$ 0.10. Further combining BAO and SNe Ia data, and relaxing the restriction on the numbers of neutrino species results in a model with $w_0$ = -1.08 $\\pm$ 0.10 and $N_{eff}$ = 4.13 $\\pm$ 0.67. These data are compatible with, but do not require the presence of an additional neutrino species.  The dominant source of systematic uncertainty in our new value of $H_0$ is the zero-point uncertainty in the Cepheid period-luminosity relation, currently limited by the small numbers of long-period Cepheids having trigonometric parallaxes.  Nevertheless, this uncertainty is now more than a factor of three smaller than the zero-point uncertainty for the {\\it Key Project}.  As outlined in F11, as part of the CHP we have also already observed Cepheids at 3.6~$\\mu$m in a sample of nearby Local Group galaxies and beyond; we are undertaking several metallicity tests of the Leavitt relation at 3.6~$\\mu$m; we are measuring a Cepheid distance to the maser galaxy, NGC 4258 at 3.6~$\\mu$m; and we are calibrating the Tully-Fisher relation at mid-infrared wavelengths. Future improvements in the Cepheid zero point will come with the launch of GAIA, and an increased sample of Milky Way Cepheids and RR Lyrae stars with accurate parallaxes, which are needed to better define the zero point of the Leavitt relation. Having a value of $H_0$ with an externally well-tested and robust total error budget of  less than 2\\% appears feasible within the next decade.   \\medskip \\medskip This work is based in part on observations made with the {\\it Spitzer} Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech. We thank the staff of the {\\it Spitzer} Science Center for the rapid processing of the data that went into this and other papers in the series. Computing resources used for this work were made possible by a grant from the Ahmanson Foundation. This research made use of the NASA/IPAC Extragalactic Database (NED).    \\vfill\\eject \\centerline{\\bf References \\rm} \\vskip 0.1cm \\vskip 0.1cm  \\par\\noindent  Allende Prieto, C., Lambert, D.~L., \\& Asplund, M.\\ 2001, \\apjl, 556, L63  \\par\\noindent  Amanullah, R., Lidman, C., Rubin, D., et al.\\ 2010, \\apj, 716, 712  \\par\\noindent Benedict, G. F., McArthur, B.E., Feast, M.W., et al.  2007, AJ, 133, 1810  \\par\\noindent Dolgov, A. D., 2002, Phys. Rep., 370, 333  \\par\\noindent Dunkley, J., Komatsu, E., Nolta, M.~R., et al.\\ 2009, \\apjs, 180, 306  \\par\\noindent  Efstathiou, G., \\& Bond, J.~R.\\ 1999, \\mnras, 304, 75  \\par\\noindent Fazio, G.G., Hora, J. L., Allen, L. E., et al. 2004, ApJS, 154, 10  \\par\\noindent Ferrarese, L., Silbermann, N.A., Mould, J.R., et al. 2000, PASP, 112, 177  \\par\\noindent Freedman, W.L., Madore, B.F., Gibson, B.K., et al. 2001, ApJ, 553, 47 (F01)  \\par\\noindent Freedman, W.L., \\& Madore, B.F. 2010, ARA\\&A, 48, 673  \\par\\noindent Freedman, W.L., \\& Madore, B.F. 2011, \\apj, 734, 46  \\par\\noindent Freedman, W.L., Madore, B.F., Scowcroft, V., et al. 2011, \\aj, 142, 192 (F11)  \\par\\noindent Hicken, M., Wood-Vasey, W.M., Blondin, S., et al.  2010, \\apj, 700, 1097  \\par\\noindent Hu, W, 2005.  ASP Conf. Series, {\\it Observing Dark Energy}, 339, 215  \\par\\noindent Komatsu, E., Dunkley, J., Nolta, M.~R., et al.,  2009, \\apjs, 180, 330  \\par\\noindent Komatsu, E., Smith, K.M., Dunkley, J., et al. 2011, ApJS, 192, 18  \\par\\noindent Lesgourgues, J., \\& Pastor, S. 2006, \\physrep, 429, 307  \\par\\noindent Lewis, A., \\& Bridle, S.\\ 2002, \\prd, 66, 103511  \\par\\noindent Ma, C.-P., \\& Bertschinger, E.\\ 1995, \\apj, 455, 7  \\par\\noindent Marconi, M., Musella, I., \\& Fiorentino, G. 2005, ApJ, 632, 590  \\par\\noindent Mehta, K.~T., Cuesta, A.~J., Xu, X.,  2012, arXiv:1202.0092  \\par\\noindent Meixner, M., Gordon, K.~D., Indebetouw, R., et al.\\ 2006, \\aj, 132, 2268  \\par\\noindent Monson, A., Freedman, W.L., Madore, B.F., et al.  2012, \\apj, submitted  \\par\\noindent Persson, S.E., Madore, B.F., Krzeminski, W., et al. 2004, AJ, 128, 2239  \\par\\noindent Reid, B.~A., Percival, W.~J., Eisenstein, D.~J., et al. 2010, MNRAS, 404, 60  \\par\\noindent Rieke, G.H., \\& Lebofsky, M.J. 1985, ApJ, 288, 618  \\par\\noindent Riess, A.~G., Strolger, L.-G., Casertano, S., et al.\\ 2007, \\apj, 659, 98  \\par\\noindent Riess, A.~G., Macri, L., Casertano, S., et al.\\ 2009, \\apj, 699, 539  \\par\\noindent Riess, A.G., Macri, L., Casertano, S., et al.  2011, ApJ, 730, 119  \\par\\noindent Romaniello, M., Primas, F., Mottini, M. et al. 2008, A\\&Ap, 488, 731  \\par\\noindent Sandage, A., Reindl, B., \\& Tammann, G.~A.\\ 2010, \\apj, 714, 1441  \\par\\noindent Scowcroft, V., Freedman, W.L., Madore, B.F.,  et al.  2011, \\apj, 743, 76  \\par\\noindent Scowcroft, V., Freedman, W.L., Madore, B.F.,  et al.  2012, \\apj, 747, 84  \\par\\noindent Stetson, P.~B.\\ 1998, \\pasp, 110, 1448  \\par\\noindent Sullivan, M., Guy, J., Conley, A., et al.\\ 2011, \\apj, 737, 102  \\par\\noindent Suyu, S. H., Treu,T., Blandford,  R. D.,  et al. 2012, arXiv:1202.4459  \\par\\noindent Turnbull, S.~J., Hudson, M.~J., Feldman, H.~A., et al.\\ 2011, \\mnras, 2088  \\par\\noindent Turner, E.L., Cen, R., \\& Ostriker, J.P. 1992, AJ, 103, 1427  \\par\\noindent Walker, A.R., 2011, Astrophys \\& Sp Sci, arXiv:1112.3171  \\par\\noindent Werner, M.W., Roellig, T. L., Low, F. J.  et al. 2004, ApJS, 154, 1  \\par\\noindent Windmark, F., Lindegren, L., \\& Hobbs, D.  2011, \\aap, 530, A76   \\vfill\\eject  \\vfill\\eject"
    },
    "1208/1208.5287_arXiv.txt": {
        "abstract": "We report polarization measurements in two prompt emissions of gamma-ray bursts, GRB~110301A and GRB~110721A, observed with the Gamma-ray burst polarimeter (GAP) aboard IKAROS solar sail mission. We detected linear polarization signals from each burst with polarization degree of $\\Pi = 70 \\pm 22$~\\% with statistical significance of $3.7~\\sigma$ for GRB~110301A, and $\\Pi = 84^{+16}_{-28}$~\\% with $3.3~\\sigma$ confidence level for GRB~110721A. We did not detect any significant change of polarization angle. These two events had shorter durations and dimmer brightness compared with GRB~100826A, which showed a significant change of polarization angle, as reported in \\citet{yonetoku2011b}. Synchrotron emission model can be consistent with all the data of the three GRBs, while photospheric quasi-thermal emission model is not favorable. We suggest that magnetic field structures in the emission region are globally-ordered fields advected from the central engine. ",
        "introduction": "Gamma-ray bursts (GRBs) are the most energetic explosions in the universe. Since the discovery of GRBs in late 1960s, spacial distribution (directions and redshifts), lightcurves and energy spectra have been observed for large amount of GRBs. However, the emission mechanism of prompt GRBs, how to release the huge energy in short time duration, is still a crucial issue. The other characteristic of electro-magnetic wave, i.e. polarization in gamma-ray band, is thought to be a key to solve the problem. Polarization measurement can probe the existence of magnetic fields, and/or some kinds of geometrical structures, which are difficult to be examined from time variability and energy spectra.  Several previous works reported marginal detections of linear polarization. \\citet{coburn2003} reported the detection of strong polarization from GRB~021206 by {\\it RHESSI}. However, independent authors analyzed the same {\\it RHESSI} data, and concluded that any polarization signals were not confirmed \\citep{rutledge2004, wigger2004}. \\citet{kalemci2007, mcglynn2007, gotz2009} reported marginal detections of polarization with $\\sim 2 \\sigma$ confidence from GRB~041219 by {\\it INTEGRAL}-SPI and -IBIS data. As the authors themselves discussed that the possibility of instrumental systematics could not be completely removed, and a part of these results were obviously inconsistent with each other \\citep[see Section~1 of][]{yonetoku2011b}. Therefore, further observations with reliable instruments are strongly required to confirm the existence of polarization in prompt GRBs.  Recently, Gamma-ray burst polarimeter \\citep[GAP;][]{yonetoku2006, yonetoku2011a, murakami2010} aboard the {\\it IKAROS} solar-power-sail \\citep{kawaguchi2008, mori2009} measured a linear polarization of $\\Pi = 27 \\pm 11$~\\% with $2.9~\\sigma$ confidence level in extremely bright GRB~100826A \\citep{yonetoku2011b}. It was found that the polarization angle changes during the prompt emission with $3.5~\\sigma$ confidence level. The GAP detector is designed for the gamma-ray polarimetry of prompt GRBs, and also realized the quite small systematic uncertainty of 1.8~\\% level \\citep{yonetoku2011a}. This may be the most convincing detection of polarization degree so far. But this is only one case, we need more sample to establish the existence of gamma-ray polarization in prompt GRBs. In this Letter, we report two more detections of gamma-ray polarization in GRB~110301A and GRB~110721A observed with {\\it IKAROS}-GAP. We analyzed all data observed by GAP until the end of 2011, and found significant polarization signals from these events. ",
        "conclusions": "\\label{discussion} \\label{sec:discussion} Major results of the GAP observations so far are simply summarized as follows: (1) there are cases with and without a significant change of polarization angle (PA), and (2) GRB~100826A, with a long duration $T \\sim 100\\;$s, has a PA change, and its polarization degree (PD) is $\\Pi = 27 \\pm 11\\%$, while GRB~110301A and GRB~110721A, with short durations $T \\sim 10\\;$s, have no PA change, and their PDs are $\\Pi \\gtrsim 30\\%$ at $2 \\sigma$ error region (see Figure~\\ref{fig3}). Here we present some theoretical implications from these observational results for synchrotron and photospheric emission models, following \\citet{yonetoku2011b}.  First we discuss synchrotron models with different types of magnetic field structure. One type is helical magnetic fields, which may be advected from the central engine through the jet. Such globally-ordered fields can produce high-PD emission \\citep{granot03,lyutikov03}. We call this `SO model'. In this model, the existence of a significant PA change implies that emission region consists of multiple patches with characteristic angular size $\\theta_p$ much smaller than the opening angle of the jet $\\theta_j$ \\citep{yonetoku2011b}. Relativistic beaming effect only allows us to observe angular size of $\\sim \\Gamma^{-1}$ around the line of sight (LOS), where $\\Gamma$ is the bulk Lorentz factor of the jet. In the case of $\\Gamma^{-1} \\sim \\theta_j$, it is natural that we see multiple patches with different magnetic field directions, and then we observe significant PA changes. On the other hand, if $\\Gamma^{-1} \\ll \\theta_j$, we only see a limited range of the curved magnetic fields, which leads to no significant PA change even for patchy emission. The net PD of the latter case can be estimated as $\\Pi \\sim 40\\%$ for photon index $\\Gamma_B \\sim -1$ (whereas synchrotron emission from a point source has $\\Pi_{\\rm max}^{\\rm syn} = (-\\Gamma_B)/(-\\Gamma_B+\\frac{2}{3})$) \\citep[see the case of $y_j = 100$ in Figure 2 of][]{toma2009}, which is compatible to the observed PD of $\\Pi \\gtrsim 30\\%$ from GRB~110301A and GRB~110721A. Therefore, the case of $\\Gamma^{-1} \\sim \\theta_j$ could correspond to GRB~100826A, while the case of $\\Gamma^{-1} \\ll \\theta_j$ to the other two bursts with no PA change.  We may consider an alternative model in which GRB jets consist of multiple impulsive shells which have globally-ordered transverse magnetic fields with different direction for each shell. Let us call this `non-steady SO model'. It has been recently claimed that such impulsive shells can be accelerated to relativistic speeds \\citep{granot12}. This polarization model may also apply to a scenario in which initial globally-ordered helical fields get distorted, making different field directions within the angular size of $\\Gamma^{-1}$, as considered in ICMART model \\citep{zhang11}. In this model, PA changes can naturally occur, although it does not necessarily occur when the number of shells is small so that $T$ is relatively small. The net PD will be $\\Pi \\sim \\Pi_{\\rm max}^{\\rm syn}/\\sqrt{N}$, where $N$ is the number of shells with different field directions. This could be as high as $\\Pi_{\\rm max}^{\\rm syn} \\sim 60\\%$ for $\\Gamma_B \\sim -1$ when $T$ is relatively small, which is consistent with the observed PD of $\\Pi \\gtrsim 30\\%$.  Shocks formed in the jet may produce sizable magnetic fields with random directions on plasma skin depth scales. Synchrotron emission from such fields can have high PD, provided that the field directions are not isotropically random but random mainly in the plane parallel to the shock plane (i.e., perpendicular to the local direction of expansion) \\citep{granot03,nakar03}. We call this `SR model'. The linear polarization directions in the observer frame are symmetric around the LOS. Thus, if the emission is patchy, net PD can be nonzero and can have PA changes \\citep{yonetoku2011b} \\citep[see also][]{lazzati09}. The characteristic angular size of the patches may be hydrodynamically constrained to be $\\theta_p \\gtrsim \\Gamma^{-1}$. In this model, the PD of emission from one patch strongly depends on the viewing angle of the patch $\\theta_v$, i.e., the angle between the axis of the patch and the LOS, and it is limited as $\\Pi \\lesssim 40\\%$ \\citep[see the cases of $y_j \\geq 1$ in Figure 3 of][]{toma2009}. This implies that fine tuning, $\\theta_v \\sim \\theta_p + \\Gamma^{-1}$, is required to have $\\Pi > 30\\%$. Furthermore, the bursts we observed are all very bright, which suggests that some patches are seen with $\\theta_v \\lesssim \\theta_p$, for which $\\Pi \\ll 40\\%$. This suggests that the SR model is not favorable to explain the observed PD of $\\Pi \\gtrsim 30\\%$.  Internal shocks may also produce strong magnetic fields with random directions on hydrodynamic scales \\citep{inoue11,gruzinov99}. We call this `SH model', which produces net PD as $\\Pi \\sim \\Pi^{\\rm syn}_{\\rm max}/\\sqrt{N}$, where $N$ is the number of independent patches with coherent magnetic field in the observable angular size $\\sim \\Gamma^{-1}$, and can naturally lead to PA changes. Unlike the SR models, the emission from patches seen with small $\\theta_v$ has high PD, so that this model is in agreement with the high brightness of the bursts. However, recent MHD simulations of internal shocks with initial density fluctuations imply $N \\sim 10^3$ \\citep{inoue11}, which cannot explain the observed PD of $\\Pi \\gtrsim 30\\%$.  Lastly, we discuss photospheric emission (Ph) model. This model assumes that the emission at $E \\gtrsim E_p$ is the quasi-thermal radiation from the photosphere. The emission at $E < E_p$ might be a superposition of many of the quasi-thermal components with different temperatures \\citep{ryde10,toma11} or contribution from synchrotron emission \\citep{vurm11}. The quasi-thermal radiation can have high PD when the radiation energy is smaller than the baryon kinetic energy at the photosphere \\citep{beloborodov11}. The linear polarization directions in the observer frame are symmetric around the LOS, same as the SR model. The PD of emission from a given point may be determined by the brightest emission, coming from just below the photosphere, which can have $\\Pi \\leq \\Pi_{\\rm max}^{\\rm qt} \\sim 40\\%$ \\citep{beloborodov11}. The PD of emission from one patch, whose angular size is hydrodynamically constrained to be $\\theta_p \\gtrsim \\Gamma^{-1}$, will reduce to $\\Pi \\lesssim 30\\%$ (in the same way as in the SR model). The patches with small viewing angle $\\theta_v$ have $\\Pi \\ll 30\\%$. Therefore, the Ph model requires very fine tuning of $\\theta_v$ to reproduce the PD of GRB~110301A, mainly at $E > E_p$ (i.e., dominated by the quasi-thermal component).  To summarize, (1) the SR, SH and Ph models are not favorable to reproduce all of our polarimetric observation results of the three GRBs, and (2) the SO and non-steady SO could explain all of them.  Recently, early optical polarization from the forward shock of a GRB afterglow has been detected as $\\Pi \\simeq 10.4 \\pm 2.5$~\\% (at $t = 149-706\\;$s) for GRB~091208B \\citep{uehara12}. If the magnetic field directions in the emission region are random on plasma skin depth scales (i.e., the SR model for a shock produced in the circumburst medium), the observed polarization reaches a maximum value around the jet break time ($\\sim 1\\;$day) \\citep{lazzati2006}. The measurement of the early polarization, higher than the typical observed late-time polarization of $\\sim 1-3$~\\% \\citep[at $t \\sim 1\\;$day;][]{covino2004}, thus disfavors this model. This result is consistent with our argument against the SR model for a shock produced within the jet.  The {\\it Fermi} satellite team has suggested that GRB 110721A has a blackbody component with temperature ranging in the $\\sim 10-100\\;$keV \\citep{axelsson12}, which almost coincide with the GAP range $70-300\\;$keV. However, their fitting model of the integrated spectrum (their Figure~2) includes the blackbody flux as only $\\sim 1/5{\\rm th}$ of the total flux in the GAP range. Thus reduction of PD by addition of the blackbody component is small. The polarization degree and angle in the GAP range is practically determined by those of the non-thermal component, which could be synchrotron emission.  The SO models assume globally-ordered fields in the emission region. On the other hand, observations have suggested that prompt emission has very high efficiency \\citep[even $>90\\%$ for some bursts;][]{zhang07,ioka06}, which means that energy dissipation, usually involving field distortion, occur globally. Reconciling high PD with high efficiency looks a dilemma, which will have to be resolved in more quantitative modeling."
    },
    "1208/1208.0377_arXiv.txt": {
        "abstract": "Earth-like planets have anelastic mantles, whereas giant planets may have anelastic cores. As for the fluid parts of a body, the tidal dissipation of such solid regions, gravitationally perturbed by a companion body, highly depends on its internal friction, and thus on its internal structure. Therefore, modelling this kind of interaction presents a high interest to provide constraints on planet interiors, whose properties are still quite uncertain.  Here, we examine the equilibrium tide in the solid central region of a planet, taking into account the presence of a fluid envelope. We first present the equations governing the problem, and show how to obtain the different Love numbers that describe its deformation. We discuss how the quality factor Q depends on the rheological parameters, and the size of the core.  Taking plausible values for the anelastic parameters, and examinig the frequency-dependence of the solid dissipation, we show how this mechanism may compete with the dissipation in fluid layers, when applied to Jupiter- and Saturn-like planets. We also discuss the case of the icy giants Uranus and Neptune. ",
        "introduction": "Once a planetary system is formed, its dynamical evolution is governed by gravitational interactions between its components, be it a star-planet or planet-satellite interaction. By converting kinetic energy into heat, the tides pertub their orbital and rotational properties, and the rate at which the system evolves depends on the physical properties of tidal dissipation. Therefore, to understand the past history and predict the fate of a binary system, one has to identify the dissipative processes that achieve this conversion of energy. Planetary systems display a large diversity of planets, with telluric planets having anelastic mantles and giant planets with possible anelastic cores \\citep{2007ARA&A..45..397U}. Since the tidal dissipation is closely related with the internal structure, one has to investigate its effects on each kind of materials that may compose a planet. Studies have been carried out on tidal effects in fluid bodies such as stars and envelopes of giant planets \\citep{2004ApJ...610..477O, 2007ApJ...661.1180O, 2009MNRAS.396..794O, 2012A&A...541A.165R}. However, the planetary solid regions, such as the mantles of Earth-like planets or the rocky cores of giant planets may also contribute to tidal dissipation \\citep[see for example][]{2012ApJ...746..150E, 2012A&A...541A.165R}. We explore here the tidal dissipation in these solid parts of planets. ",
        "conclusions": "Our evaluations reveal a much higher dissipation in the solid cores of planets than that found by \\cite{2004ApJ...610..477O} for the fluid envelope of a planet having a small solid core. These results seem to be in good agreement with observed properties of Jupiter's and Saturn's system \\citep{2009Natur.459..957L,2012ApJ...752...14L}. In the case of the ice giants Uranus and Neptune, too much uncertainties remain on internal structure to give an order of magnitude, other than a minimum value, of tidal dissipation in the solid regions, which constitutes a first step in the tudy of such planets."
    },
    "1208/1208.0007_arXiv.txt": {
        "abstract": "We perform a local (short-wavelength) linear stability analysis of an axisymmetric column of magnetized plasma with a nearly toroidal magnetic field and a smooth poloidal velocity shear by perturbing the equations of relativistic magnetohydrodynamics. We identify two types of unstable modes, which we call `exponential' and `overstable', respectively. The exponential modes are present in the static equilibria and their growth rates decrease with increasing velocity shear. The overstable modes are driven by the effects of velocity shear and dominate the exponential modes for sufficiently high shear values. We argue that these local instabilities can provide an important energy dissipation mechanism in astrophysical relativistic jets. Strong co-moving velocity shear arises naturally in the magnetic acceleration mechanism, therefore it may play a crucial role in converting Poynting-flux-dominated jets into matter-dominated jets, regulating the global acceleration and collimation processes, and producing the observed emission of blazars and gamma-ray bursts. ",
        "introduction": "Many high-energy astrophysical phenomena are powered by jets that form in the vicinity of compact objects and are accelerated to relativistic speeds (blazars and gamma-ray bursts (GRBs)) or very high subrelativistic speeds (X-ray binaries, pulsar wind nebulae). The currently standard model of magnetic acceleration requires that jets be initially dominated strongly by their Poynting flux, and that acceleration takes a few orders of magnitude in distance scale \\citep[\\eg,][]{2007MNRAS.380...51K,2009ApJ...698.1570L}. In this process, the jet expands laterally far beyond the light cylinder of the central object, stretching the magnetic field into a nearly toroidal configuration. Magnetically-dominated plasma with a toroidal magnetic field is known to be highly susceptible to Current-Driven Instability (CDI). These instabilities cause severe problems in laboratory experiments on `Z-pinch' plasma systems. However, in the context of astrophysical jets they provide both a challenge and an opportunity. The challenge is not to destroy the jet as it propagates over several orders of magnitude in distance. The opportunity is to provide a mechanism of energy dissipation --- by tapping free magnetic energy and converting it to particle energy --- that can power high-energy emission from blazars and GRBs \\citep{2006A&A...450..887G}.  Global, three-dimensional general-relativistic magnetohydrodynamical (3D GRMHD) numerical simulations of jets, tracing them from their formation around spinning black holes over distances of thousands of gravitational radii, indicate that they are marginally stable to long-wavelength modes of CDI, including the most dangerous pinch ($m=0$) and kink ($m=1$) modes \\citep{2009MNRAS.394L.126M}. On the other hand, it is expected that the growth rate of CDI is similar to the Alfv{\\'e}n crossing time scale of the jet, and thus should not be sensitive to the wavelength. Short-wavelength (local) modes are particularly important for the problem of energy dissipation, since the dissipation rate scales like the inverse square of the wavelength. These modes were predicted analytically to play an important role in AGN jets and pulsar wind nebulae \\citep{1998ApJ...493..291B}, and recent local numerical simulations have confirmed their general properties \\citep{2009ApJ...700..684M,2011ApJ...728...90M,2012MNRAS.422.1436O}.  Most analytical studies of CDI have been performed in the force-free approximation, i.e. neglecting the pressure and inertia of the matter content of jet plasma \\citep{1996MNRAS.281....1I,1999MNRAS.308.1006L,2001PhRvD..64l3003T,2009ApJ...697.1681N}. This approximation is valid in the innermost regions of jets, within the fast magnetosonic point, at which the plasma magnetization $\\sigma$ is still substantial. Beyond the fast magnetosonic point, the magnetic nozzle effect results in further gradual jet acceleration, and a decrease in the jet magnetization to the order of unity \\citep{1994ApJ...426..269B}. In this jet region, full MHD force balance is necessary to study plasma stability, as was done by \\cite{1998ApJ...493..291B}. However, that study assumed a static background configuration, while it is known that the efficiency of jet acceleration is strong function of cylindrical radius \\citep{2007MNRAS.380...51K,2008MNRAS.388..551T}. Within the jet acceleration region, a strong shear of the poloidal velocity can be expected to arise, and this effect needs to be incorporated in the local stability analysis.  Velocity shear is the driving factor of the Kelvin-Helmholtz Instability (KHI), which has been studied extensively in the context of astrophysical jets. Most analytical studies of KHI have adopted a vortex-sheet background configuration, i.e. a parallel discontinuity of the velocity field \\citep{1976MNRAS.176..421T,1976MNRAS.176..443B,1978A&A....64...43F,1984MNRAS.208..887B,2007ApJ...664...26H}. In some works, a more gradual transition region of finite width was considered \\citep{1982MNRAS.198..617R,1982MNRAS.198.1065F,1991MNRAS.252..505B}. However, we have found no analytical study that deals with a strictly smooth velocity shear, which is natural for the local stability problem that we study in the context of CDI. Several numerical studies have investigated the effect of velocity shear on CDI or the interaction between CDI and KHI modes. Some of them have adopted the vortex-sheet approximation \\citep{2011ApJ...734...19M}, while some use a gradual velocity shear with a well-localized region of strong shear generating typical KHI modes \\citep{2002ApJ...580..800B}. For moderate shear values, the distinction between CDI and KHI modes becomes quite vague.  In this work we study local CDI modes of cylindrical MHD equilibria in the presence of strong but gradual velocity shear. In Section \\ref{sec_disp}, we derive the dispersion relation for a sheared plasma column, which is given by Equation (\\ref{eq_disp}). In Section \\ref{sec_stab}, we explore the complex space of mode frequencies, looking for the fastest-growing mode as a function of the background equilibrium parameters. In Section \\ref{sec_appl}, we discuss the relevance of our results to global models of astrophysical jets. We summarize our results in Section \\ref{sec_summ}.    In this work, we use cgs units with $c=1$. In the linear analysis, all quantities describing the background configuration are denoted with subscript `0' (zero-order terms), while all quantities describing the linear perturbation are denoted with subscript `1' (linear terms). Vector quantities are denoted with a bold font. ",
        "conclusions": "\\label{sec_appl}  In order to evaluate the importance of the velocity shear in affecting the stability of astrophysical jets, we first need to estimate the magnitude of this parameter. Our calculations in previous sections were performed in the local co-moving frame (in this section, we denote the quantities calculated in the co-moving frame with a prime). The co-moving velocity shear is directly related to the transverse gradient of the Lorentz factor calculated in the external frame: \\be \\label{eq_u_Gamma} u'=r\\partial_rv_{\\rm 0,z}'=\\Gamma^2r\\partial_rv_{\\rm 0,z}=\\frac{1}{v_{\\rm 0,z}}\\pder{\\ln\\Gamma}{\\ln r}\\,. \\ee In relativistic jets, a co-moving velocity shear value of $u'=1$ corresponds to the doubling of the Lorentz factor (as measured by an outside observer) with a doubling of the radius. The co-moving velocity shear is significant when the Lorentz factor varies by a sizable factor across the jet.  Relativistic jets are thought to be accelerated by the action of magnetic stresses in an extended region dominated by the Poynting flux. Both analytical studies \\citep{2004ApJ...605..656V,2009ApJ...698.1570L} and numerical simulations \\citep{2007MNRAS.380...51K,2008MNRAS.388..551T,2010ApJ...709.1100P} show that this process does not operate with uniform efficiency across the jet radius. This results in a strong dependence of the Lorentz factor within the jet on the radial coordinate (equivalently on the polar angle). To illustrate these dependencies, in the top panel of Figure \\ref{fig_toy1}, we plot three models for the radial variation of the Lorentz factor, qualitatively based on the results of \\cite{2007MNRAS.380...51K} (models C1 and C2) and \\cite{2008MNRAS.388..551T} (Appendix 4.1). In the bottom panel of Figure \\ref{fig_toy1}, we plot the corresponding values of the co-moving velocity shear given by Equation (\\ref{eq_u_Gamma}). We find that velocity shear values of order $\\left|u'\\right|\\sim 0.5$ and higher can be expected in relativistic jets. This indicates that the effect of the poloidal velocity shear on the stability of jets is very important. A gradual jet acceleration also means that the importance of velocity shear can be expected to increase with distance.   \\begin{figure} \\includegraphics[width=\\columnwidth]{figs/toy1.eps} \\caption{\\emph{Top panel:} three functions for the Lorentz factor $\\Gamma$ distribution across a jet, motivated by numerical results of Komissarov et~al. (2007) and Tchekhovskoy et~al. (2008). \\emph{Bottom panel:} the corresponding distribution of the co-moving velocity shear $u'$.} \\label{fig_toy1} \\end{figure}   In our local approximation, the sign of the velocity shear has no effect on the maximum growth rate for a given system; only the magnitude of the shear is important. This is because in Equation (\\ref{eq_disp}), the linear term in $u$ is coupled with the $l/k$ ratio, and a change in the sign of $u$ can be countered by a change in the sign of $l$ (inward or outward propagation) or $k$ (upstream or downstream propagation). On the other hand, \\cite{2009ApJ...697.1681N} found that for global modes in the force-free approximation, the sign of the poloidal velocity shear is important. They found that jets with positive velocity shear (Lorentz factor increasing with radius) are stable, while those in which the velocity shear changes sign at an intermediate radius, as in the example shown in our Figure \\ref{fig_toy1} with a dotted line, are unstable. This fundamental difference in conclusions reached via global and local analysis is not surprising. It demonstrates that even in globally stable jets, vigorous dissipation may result from the presence of local modes.  The local stability criterion formulated by \\cite{1998ApJ...493..291B} for a static plasma column, which can be written as $a>2(\\alpha_{\\rm\\phi}+1)$, is no longer valid if poloidal velocity shear is present. This is clearly illustrated in Figure \\ref{fig_disp1_mix2}. In the case of $\\alpha_{\\rm\\phi}=-0.5$, the stability criterion is $a>1$. However, even for a very small value of $u$, there exist overstable modes with $a>1$. Even for $\\alpha_{\\rm\\phi}\\le -1$, in which case the static equilibria are always stable, we find overstable solutions.  A mode with a co-moving dimensionless growth rate of $\\tilde\\omega_{\\rm I}'$ will double its amplitude over the co-moving time scale of $t_2'=0.7r/\\tilde\\omega_{\\rm I}'$. Over this time, an element located initially at $z$ will travel a distance $z_2$ in the external frame given by $(z_2/z)\\sim 0.7(\\Gamma\\theta)/\\tilde\\omega_{\\rm I}'$, where $\\Gamma\\gg 1$ is the local Lorentz factor and $\\theta=r/z$ is the polar angle. Since in relativistic jets we expect $\\Gamma\\theta\\lesssim 1$, for $\\tilde\\omega_{\\rm I}'\\sim\\mathcal{O}(1)$ the growth length measured in the external frame is not particularly short, with $(z_2/z)\\sim\\mathcal{O}(1)$, with the exception of the core region close to the jet axis. For a typical co-moving growth rate of $\\tilde\\omega_{\\rm I}'\\sim 0.4$, the energy dissipation associated with the local CDI can be expected to proceed over at least one order of magnitude in distance along the jet. A major caveat is that the growth of local CDI modes competes with lateral jet expansion. The question of whether jet expansion is able to quench the CDI modes must be addressed by a careful analysis of realistic jet models.  The results of this work need to be confirmed by numerical simulations. We are currently preparing local simulations of cylindrical plasma columns similar to those in \\cite{2012MNRAS.422.1436O}, adding a background poloidal velocity shear. However, this issue also needs to be addressed by global simulations, like those in \\cite{2009MNRAS.394L.126M}, but with a much higher resolution.  Our study of the local development of CDI is still very simplified, as we do not take into account several effects that could potentially have an impact on the CDI growth rate. Some of these effects, like rotation or lateral expansion, are not supposed to be important at the scales we are interested in. However, centrifugal forces due to jet collimation, enhanced by relativistic effects, can significantly change the radial force balance \\citep{2009ApJ...697.1681N} and should be included in future work."
    },
    "1208/1208.2528_arXiv.txt": {
        "abstract": "{There are only about 65 R Coronae Borealis stars known in our Galaxy, and none in globular clusters. As these stars are thought to result from the merger of two white dwarfs, one would expect the higher stellar density of globular clusters to favor their formation.  We have searched for such stars in Galactic globular clusters, as their presence in a specific category of clusters might provide more clues as to their formation.  We selected from the WISE all-Sky source catalog all the stars within the tidal radius of the 150 globular clusters within 50 kpc, which is the distance to which RCB stars are detectable by WISE. The total number of stars selected in this way was 635989. We then successively applied the eight selection criteria of Tisserand (2012) satisfied by RCB stars to the dereddened photometric WISE and 2MASS data.  Only three stars satisfying the conditions were found in the field of three globular clusters. The star in the field of Liller 1 is most probably a protostar. For the two other candidates, the absence of photometry in the visible range did not allow us to establish their nature with certainty. We further identified one dust-enshrouded star that only satisfied the first selection criteria, and used DUSTY to determine that it is a star of temperature 4800K enshrouded in a dusty envelope with a temperature 300 K and an opacity in the visible of 0.59. It is probably an Xray binary star with a dusty accretion disk.  We found no RCB stars truly belonging to a globular cluster, thus providing a constraint on their formation mechanism. } ",
        "introduction": "\\label{intro}  R Coronae Borealis stars are a very small group of hydrogen-deficient and carbon-rich supergiants. About 65 are known in the Galaxy and 25 in the Magellanic Clouds. In our Galaxy, they are mostly confined to low Galactic latitudes, and are thus thought to be part of the bulge population. Two scenarios have been proposed to explain their existence, either the double degenerate merger of two white dwarfs, or the final helium shell flash of the central star in a planetary nebula (see Clayton 2012 for a review).  The expected number of RCB stars in the Galaxy can be as large as 500 (Tisserand 2012) or even 5000 (Clayton 2012).  It thus seems feasible to increase the number of known RCB stars significantly, in order to better understand their evolutionary status. Since these stars are carbon-rich, they should be more easily detectable from their infrared properties.  Tisserand (2012) recently proposed eight criteria for identifying RCB stars from their mid-infrared WISE colors.  Our own interest lies with the stellar populations in globular clusters (Sharina \\& Davoust 2009, Sharina et al. 2010), and we recently discovered a CH star in the globular cluster NGC 6426 (Sharina et al. 2012). Since the presently more likely formation scenario for RCB stars seems to be a merger of white dwarfs, the dense environment of globular clusters should be ideal for their formation. We thus decided to take advantage of the recently available WISE all-sky source catalogue to search for RCB stars in globular clusters, using the selection criteria of Tisserand (2012). An advantage of searching for such stars in globular clusters is that the distance and reddening are known {\\it a priori}. ",
        "conclusions": "\\label{conc}  The negative result of our search for RCB stars in Galactic globular clusters indicates that such stars are very rare, if not completely absent, in these old stellar systems. This may be because the RCB phenomenon has a very short lifetime, is a rare event in the evolution of stars, or that it occurs in stars younger than five to seven Gyr, which is the age range of the youngest clusters explored in this project."
    },
    "1208/1208.0231_arXiv.txt": {
        "abstract": "{Although numerous archival \\xmm\\ observations existed towards the Small Magellanic Cloud (SMC) before 2009, only a fraction of the whole galaxy had been covered.} {Between May 2009 and March 2010, we carried out an \\xmm\\ survey of the SMC, to ensure a complete coverage of both its bar and wing. Thirty-three observations of 30 different fields with a total exposure of about one Ms filled the previously missing parts.} {We systematically processed all available SMC data from the European Photon Imaging Camera. After rejecting observations with very high background, we included 53 archival and the 33 survey observations. We produced images in five different energy bands. We applied astrometric boresight corrections using secure identifications of X-ray sources and combined all the images to produce a mosaic covering the main body of the SMC.} {We present an overview of the \\xmm\\ observations, describe their analysis, and summarise our first results, which will be presented in detail in follow-up papers. Here, we mainly focus on extended X-ray sources, such as supernova remnants (SNRs) and clusters of galaxies, that are seen in our X-ray images.} {Our \\xmm\\ survey represents the deepest complete survey of the SMC in the 0.15--12.0 keV X-ray band. We propose three new SNRs that have low surface brightnesses of a few \\oergcm{-14} arcmin$^{-2}$ and large extents. In addition, several known remnants appear larger than previously measured at either X-rays or other wavelengths extending the size distribution of SMC SNRs to larger values.} ",
        "introduction": "The study of X-ray source populations and diffuse X-ray emission in nearby galaxies is of major importance to improve our understanding the X-ray output of more distant galaxies as well as learning about processes that occur on interstellar scales within our own Galaxy. \\xmm\\ and \\chandra\\ were used to perform deep X-ray surveys of the Local Group galaxies M31 \\citep{2005A&A...434..483P,2011A&A...534A..55S} % and M33 \\citep{2004A&A...426...11P,2006A&A...448.1247M,2011ApJS..193...31T}. % These deep observations of M31 and M33 allow for the first time the study of large samples of different source classes (with limiting point-source luminosities of $\\sim$\\oergs{35}) in a galaxy other than the Milky Way. This includes, e.g., the study of about 80 supernova remnant (SNR) candidates in M33, which is the largest sample of remnants detected at optical and X-ray wavelengths in any galaxy \\citep{2010ApJS..187..495L}, % optical novae as the major class of supersoft X-ray sources (SSSs) in M31 and M33 \\citep{2005A&A...442..879P,2007A&A...465..375P,2010A&A...523A..89H,2011A&A...533A..52H} % and the discovery of type-I X-ray bursts from neutron star X-ray binaries in M31 globular clusters \\citep{2005A&A...430L..45P}. %  The Large and Small Magellanic Clouds (LMC and SMC), nearby neighbours of the Milky Way, have different chemical compositions of low metallicity, are irregular in shape, and are strongly interacting both with the Milky Way and each other. These properties influence their star formation history and therefore, any study of stellar populations in the Magellanic Clouds (MCs) is particularly rewarding \\citep[see e.g. ][ with respect to the Be/X-ray binary population of the SMC]{2010ApJ...716L.140A}. Their proximity makes them ideal targets for X-ray studies. Limiting point-source luminosities of a few \\oergs{33} (a factor of $\\sim$50 lower than for M31 and M33) are reached with \\xmm\\ and extended objects like supernova remnants can easily be resolved \\citep[at the SMC distance of 60 kpc, ][ the angular resolution of $\\sim$10\\arcsec\\ provided by \\xmm\\ corresponds to a linear size of 3 pc]{2005MNRAS.357..304H}.  Previous X-ray surveys of the MCs performed with the imaging instruments of the \\einstein\\ \\citep{1981ApJ...248..925L,1991ApJ...374..475W,1992ApJS...78..391W}, % ASCA \\citep[only SMC; ][]{2003PASJ...55..161Y} % and ROSAT \\citep{1999A&AS..136...81K,2000A&AS..142...41H,2000A&AS..147...75S,1999A&AS..139..277H,2000A&AS..143..391S} satellites revealed discrete X-ray sources and large-scale diffuse emission. In particular, the high sensitivity and the large field of view of the ROSAT PSPC provided the most comprehensive catalogues of discrete X-ray sources ever compiled and revealed the existence of a hot thin plasma in the interstellar medium (ISM) of the MCs with temperatures between 10$^6$ and 10$^7$ K \\citep{2002A&A...392..103S}. %  However, owing to their relatively small distance, the MCs extend over large areas on the sky and high spatial resolution X-ray surveys with modern instrumentation need to complete a large number of raster observations. \\chandra\\ and \\xmm\\ observed various targets in the MCs and only for the SMC do these `serendipitous' surveys cover a significant part of the galaxy. The disconnected fields of the \\chandra\\ Wing Survey \\citep{2008MNRAS.383..330M} are distributed around the eastern wing. The observations available in the \\xmm\\ archive at the beginning of 2009 mainly covered parts of the SMC bar and wing with very different exposure times. To fill the gaps between the archival observations, we successfully applied for \\xmm\\ observations of thirty fields in the SMC, which were performed between May 2009 and March 2010. Together with the archival data, the new observations cover the full extent of the SMC with the European Photon Imaging Camera \\citep[EPIC, ][]{2001A&A...365L..18S,2001A&A...365L..27T} on board \\xmm.  First studies using the new \\xmm\\ survey data mainly focussed on individual objects belonging to the SMC. Two known SSSs \\citep{1996LNP...472Q.299G} were observed in bright state during our \\xmm\\ survey. SMP SMC 22 \\citep{1978PASP...90..621S} is the most X-ray luminous central star of a planetary nebula in the SMC \\citep{2010A&A...519A..42M} and the symbiotic binary SMC\\,3 is a highly variable X-ray source that was seen at its highest intensity state observed so far during our survey \\citep{2011A&A...529A.152S}. A faint SSS identified with a Be star could be the first Be/white dwarf system discovered in the SMC \\citep{2012A&A...537A..76S}. New Be/X-ray binaries were discovered in outburst during the \\xmm\\ survey observations. For two of them, X-ray pulsations were discovered in the EPIC data \\citep{2011A&A...527A.131S,2011MNRAS.414.3281C}, while two new transients show all the characteristics of Be/X-ray binaries, but no significant pulsations were detected \\citep{2012MNRAS.424..282C}. Another new Be/X-ray binary was discovered in a search for highly absorbed X-ray binaries in the EPIC data \\citep{2011A&A...532A.153N}. \\citet{2011A&A...530A.132O} combined all available \\xmm\\ data of the SNR IKT\\,16 to study its morphology and X-ray spectrum.  Here, we present the first EPIC mosaic images from the survey and focus on extended X-ray sources such as SNRs in the SMC and galaxy clusters in the background. A detailed analysis of the images concerning the detection of point sources is presented in \\defcitealias{Sturm_cat}{SHP12}\\citet[][ hereafter \\citetalias{Sturm_cat}]{Sturm_cat}, which resulted in 5236 detections of 3053 individual sources. ",
        "conclusions": "We have performed a deep X-ray survey of the SMC using the EPIC instruments on board \\xmm\\ between May 2009 and March 2010, which, together with archival observations, covers the bar and eastern wing of the galaxy. While a detailed analysis of the population of discrete X-ray sources (unresolved and with small extents) is described in \\citetalias{Sturm_cat}, we present here first mosaic images of the main body of the SMC. We study the morphology of more extended X-ray emission from supernova remnants and galaxy clusters. Several galaxy clusters can easily be identified in our EPIC images by their extended appearance and X-ray colours.  We have found that several of the known SNRs listed with their spatial extent in the catalogue of \\citetalias{2010MNRAS.407.1301B} appear larger in the X-ray images. This is particularly true for large SNRs with low temperatures (suggesting old remnants) and low surface brightnesses, which can be better mapped with the high sensitivity of \\xmm\\ in soft X-rays. However, a larger extent could also be caused by the presence of more than one SNR merging into larger bubbles. In these cases, the consistent X-ray colours require the neighbouring remnants, which could have formed in star clusters, to have similar ages.  We propose three new SNRs with large extents, low surface brightnesses, and soft X-ray spectra (labelled with their XMM names in Fig.~\\ref{fig-snrima}). The high sensitivity of the EPIC instruments at low energies allowed us to detect their faint and soft emission with a typical surface brightness of \\oergcm{-14} arcmin$^{-1}$ (0.2-2.0 keV). Only one of the three remnants (XMMU\\,J0056.5$-$7208) is seen in the MCELS images in the \\Halpha, [\\SII], and [\\OIII] emission lines. Shell-like optical filaments suggest that we see either one SNR extending further to the south-west of the bulk of the X-ray emission or two merging SNR shells with the southern one being fainter in X-rays but brighter in the optical than the northern shell. Our survey covers all other known SNRs in the SMC, except the recently discovered faint SNR far out in the eastern wing, which is located outside our survey area \\citep{2012A&A...537L...1H} and also exhibits a low surface brightness and low temperature, indicative of older SNRs. For completeness, the image of this remnant (designated XMMU\\,J0127.7$-$7332), which harbours the Be/X-ray binary pulsar SXP\\,1062 at its centre is included in the image gallery of Fig.~\\ref{fig-snrima}. Overall, our studies of the X-ray morphology of the SNRs in the SMC extend their size distribution to larger values. In particular for B0050$-$728, the largest SNR known in the SMC, we have found an X-ray diameter as large as 7', which corresponds to 122 pc at a distance of 60 kpc. In the MCs, this is only exceeded by the LMC SNR DEM\\,L203 \\citepalias{2010MNRAS.407.1301B}."
    },
    "1208/1208.4624_arXiv.txt": {
        "abstract": "We revisit the early evolution of the Moon's bombardment.  Our work combines modeling (based on plausible projectile sources and their dynamical decay rates) with constraints from the lunar crater record, radiometric ages of the youngest lunar basins, and the abundance of highly siderophile elements in the lunar crust and mantle. We deduce that the evolution of the impact flux did not decline exponentially over the first billion years of lunar history, but also there was no prominent and \u201cnarrow\u201d impact spike $\\sim 3.9$ Gy ago, unlike that typically envisioned in the lunar cataclysm scenario.  Instead, we show the timeline of the lunar bombardment has a sawtooth-like profile, with an uptick in the impact flux near $\\sim 4.1$ Gy ago. The impact flux at the beginning of this weaker cataclysm was 5-10 times higher than the immediately preceding period.  The Nectaris basin should have been one of the first basins formed at the sawtooth.  We predict the bombardment rate since $\\sim 4.1$~Gy ago declined slowly and adhered relatively close to classic crater chronology models (Neukum and Ivanov (1994)).  Overall we expect that the sawtooth event accounted for about 1/4 of the total bombardment suffered by the Moon since its formation.  Consequently, considering that $\\sim 12$-14 basins formed during the sawtooth event, we expect that the net number of basins formed on the Moon was $\\sim 45$-50. From our expected bombardment timeline, we derived a new and improved lunar chronology suitable for use on Pre-Nectarian surface units. According to this chronology, a significant portion of the oldest lunar cratered terrains has an age of 4.38-4.42 Gyr. Moreover, the largest lunar basin, South Pole Aitken, is older than 4.3Gy, and therefore was not produced during the lunar cataclysm. ",
        "introduction": "The temporal evolution of lunar bombardment is a subject of intense debate.  A natural expectation is that it declined with time during the early epochs of solar system history, while planetesimals left over from planet accretion were in the process of being gradually removed by dynamical and collisional mechanisms.  In this respect, a surprise came with the first analysis of the lunar samples collected by the Apollo missions.  They revealed a clustering of radiometric impact ages at about 3.9 Gy ago (Papanastassiou and Wasserburg, 1971a, 1971b; Wasserburg and Papanastassiou, 1971; Turner et al., 1973). Tera et al. (1974) concluded that a major bombardment episode occurred on the the Moon at that time, i.e. about 0.6 Gyr after the Moon formation (4.5 Gy ago; see e.g. Kleine et al., 2009), which they named {\\it terminal lunar cataclysm}.  {More recently, laboratory analyses on lunar meteorites, which should be more representative of the entire lunar surface than the Apollo samples, confirmed the strong deficit of impact ages older than $\\sim 4$~Gy (Cohen et al. 2000), although they did not show a narrow impact spike.  This absence of older ages is consistent with the cataclysm hypothesis, but it has been argued that it could also be the result of biases that work against finding samples with the oldest impact ages (Hartmann 1975, 2003; see Hartmann et al., 2000 and Chapman et al., 2007 for reviews of the contrasting arguments).}   An analysis of lunar crater densities also fails to yield an unambiguous view of the temporal evolution of lunar bombardment. Neukum and Wilhelms (1982) (see also Neukum, 1983; Neukum and Ivanov, 1994, hereafter NI94) studied the crater density over terrains of ``known'' radiometric age. {They concluded that the bombardment rate was roughly constant (within a factor of 2 or so) until 3.5 Gy ago, a result that is generally accepted today. In addition, they argued for a long {\\it smooth} decay of the impactor flux at older times}. Thus, in their model, there is no lunar cataclysm. The problem, however, is that only the youngest units, starting with the Imbrium basin ~3.8-3.9 Gy ago, have well established radiometric ages, whereas the ages of older basins, like Nectaris, are uncertain (e.g. Norman et al., 2010).  Neukum and collaborators assumed the age of Nectaris basin was $\\sim$4.1 Gy because this age appears in the samples collected by the Apollo 16 mission that landed in the lunar highlands near Nectaris (e.g., Maurer et al., 1978). In this case, the density of craters as a function of age between 4.1 to 3.5 Gy ago seems to decline as $\\exp (-a t)$, where $a=6.95$ and $t$ is measured in Gy (see Fig. 1). This exponential evolution was then extrapolated backwards in time {by NI94}, to estimate the impact flux during the oldest lunar epochs\\footnote{NI94 also used as a data-point the density of craters on the lunar highlands, but the age of the latter (which NI94 assumed to be 4.35Gy) is not well contrained.}.  A different view is summarized by Ryder (1990), St{\\\"o}ffler and Ryder (2001) and Ryder (2002).  They argued that the age of Nectaris is 3.9 Gy because this age appears more prominently than the 4.1 Gy age among Apollo 16 Descartes terrain samples. If one assumes this age, then the {\\it same} crater counts on Nectaris imply a bombardment rate that has a {\\it much steeper} decline over the 3.5--3.9 Gy period than in NI94. This steeper decline cannot be extrapolated in time back to the lunar formation event because it would lead to {unrealistic physical} implications. For instance, the Moon would have accreted more than a lunar mass since its formation (Ryder, 2002)!  Consequently this scenario implies that the bombardment rate could not have declined smoothly, but rather should have been smaller before 3.9 Gy ago than in the 3.5-3.9 Gy period, in agreement with the cataclysmic impact spike hypothesis.  In fact, in an end-member version, Ryder (1990) also suggested that {\\it all} impact basins could have formed during such a cataclysm impact spike.  The most recent analysis of Apollo 16 samples suggest that the younger ages from the Descartes terrain are probably ejecta from Imbrium (Norman et al., 2010). The older ages, however, have no diagnostic link to Nectaris basin ejecta. This means the age of Nectaris remains uncertain, and may or may not be represented among Apollo 16 samples. Thus, no definitive conclusion can be derived in favor of the cataclysm or the smooth exponential decline hypothesis from these data.  {Other studies on the lunar crater record reported support for a lunar cataclysm.  Strom et al. (2005) detected a change in the Size Frequency Distribution (SFD) of old craters (i.e. on the highlands) relative to young craters (i.e. on the maria plains). Marchi et al. (2012) detected the signature of a change in the velocity of the projectile populations hitting the Moon at Nectarian and pre-Nectarian times respectively. Both findings suggest drastic changes in the impactor populations of the solar system, consistent with the cataclysm hypothesis.  However, an opposing viewpoint has been suggested by Fassett et al. (2012). They also found two populations of projectiles, but the transition from one to the other occurred in mid-Nectarian epoch, i.e. in the middle of the putative cataclysm. They interpreted this result as problematic for the lunar cataclysm scenario.  Therefore, it is fair to say that interpreting the early cratering record of the Moon is challenging.}  In this paper, we choose not to enter into those technical debates, but instead revisit the problem with a new combination of theoretical considerations (by looking at the dynamical evolution of plausible projectile sources) and existing physical constraints.  More precisely, we look to calibrate the ``free parameters'' of the problem (i.e., size of the projectile population, timing of the instability that released the projectiles from a formerly stable reservoir, the approximate age of Nectaris basin) to produce a model that is consistent with (i) the possible dynamical evolution of the solar system, (ii) the lunar crater record and (iii) particular geochemical constraints derived from lunar samples. As we will show, our results support a view that is somewhat intermediate between the two end-member camps described above: all lunar basins forming in a smooth decline or a prominent and narrow impact spike 3.9~Gy ago.  In fact, we will argue for the need of a sudden increase in the lunar impact rate, but as early as $\\sim$4.1-4.2 Gy ago, and not one as pronounced as in Ryder's (1990) description of the lunar cataclysm. {Our view was proposed before (e.g. Fig. 3 in Hartmann et al., 2000), but never quantified through a calibrated model.}  Consequently, we believe the lunar cataclysm implies a decline of the bombardment rate since 4.1 Gy ago, in agreement with that described by NI94. ",
        "conclusions": "{We have taken our best models of the early solar system evolution, determined the impact flux on the Moon over time and then calibrated these results using the existing dynamical, geochemical, and crater density constraints.  We infer that the evolution of the lunar bombardment rate} is somewhat intermediate between the two end-member views in this historical controversy.  Our model bombardment rate from 4.1 to 3.5 Gy ago agrees with the exponential decay illustrated by NI94, the champions of the no-cataclysm view. We find that it is impossible, however, to extrapolate their exponential flux backward in time before 4.1 or possibly 4.2 Gy. We believe a discontinuity in the evolution of the bombardment rate, or a lunar cataclysm, is the easiest way to match constraints.  The timeline of the Moon's bombardment that emerges from our study has a sawtooth profile, with a moderate uptick at 4.1--4.2 Gy (see Fig 3, left panel).  This stands in sharp contrast with the prominent impact spike usually shown in sketches of the lunar cataclysm; {instead it is in broad agreement with the scenario of ``weak cataclysm'' promoted in Fig. 3 of Hartmann et al. (2000).}  Our impact flux model predicts the lunar bombardment from 4.1 Gy ago --which includes the cataclysm caused by the destabilization of the main asteroid belt ($\\sim 15-25$\\% of cataclysm impactors) and E-belt ($\\sim 75-85$\\% of cataclysm impactors)-- accounts for approximately 25\\% of the total bombardment suffered by the Moon since it formed. This is in agreement with the total number of basins on the Moon ($\\sim 50$), of which only $\\sim 12$-14 are Nectarian and early-Imbrian (i.e. younger than $\\sim 4.1$ Gy).  We note that we could be underestimating the number of basins because some ancient basins have probably been erased (Frey and Romine, 2011). The best estimate at present is that the total number of basins probably does not exceed $\\sim 60$ (Frey, private communication), not far from our expectations for the combined leftover planetesimal and E-belt models. We stress that basin formation is necessarily a stochastic process which can deviate from probabilistic expectations due to small number statistics.  {Assuming Poisson statistics (i.e. the error bar on $N$ events is $\\sqrt{N}$), the number of basins formed in our model since the destabilization of the E-belt is $12 \\pm 4$, while the total number of basins is $45 \\pm 7$.}  Large portions of the lunar highlands have a crater density that is about twice that of Nectaris (Strom, 1977; Marchi et al., 2012), with a value of $N_{20}=1.73\\times 10^-4$. According to the cumulative bombardment shown in Fig. 3, this would imply that these portions of highlands started to retain craters about 4.38-4.42 Gy ago, ages that are consistent with recent estimates of the timescale for the thickening of the lunar lithosphere (Meyer et al. 2010). This age also approximates the closure age of the crust, as derived from zircons that crystallized in the remaining urKREEP residue of the lunar magma ocean 4.38 to 4.48 Gy (Nemchin et al. 2009 ; Taylor et al. 2009). {Similarly, the value of $N_{20}$ for the SPA floor is $1.36\\times 10^{-4}$ (Marchi et al., 2012). This implies that craters started to accumulate on SPA since 4.33-4.39~Gy ago. This age should be considered as a lower bound for the formation age of SPA, because the basin's floor might have solidified only after some time; it clearly shows that SPA is an old basin, which definitely predates the cataclysm event.}  The sawtooth-like bombardment timeline has important implications for Earth's habitability.  In the no-cataclysm view, the Earth was increasingly hostile to life going back in time, as the bombardment exponentially increased.  In the classic view of the lunar cataclysm, the prominent impact spike 3.8-3.9 Gy ago conceivably sterilized the Earth by vaporizing all the oceans and thereby creating a steam atmosphere (Maher and Stevenson, 1988; however see Abramov and Mojzsis, 2009). In our sawtooth view, big impactors hit over an extended period, with more lulls and therefore more opportunities for the Hadean-era biosphere to recover. Perhaps in this scenario, life formed very early and has survived in one form or another through the lunar cataclysm.  \\vskip 20pt {\\bf Acknowledgments}  We thank R. Walker, W. Hartmann and an anonymous reviewer for the comments provided to the first version of this manuscript. A. Morbidelli and S. Marchi thank Germany\u2019s Helmholtz Alliance for providing support through their \u201cPlanetary Evolution and Life\u201d programme.  The contributions of S.Marchi, W. F. Bottke, and D. A. Kring were supported by NASA's Lunar Science Institute (Center for Lunar Origin and Evolution at the Southwest Research Institute in Boulder, CO; Center for Lunar Science and Exploration at the Lunar and Planetary Institute in Houston, TX)."
    },
    "1208/1208.3697_arXiv.txt": {
        "abstract": "Current evidence for dark matter in the universe does not exclude heavy composite nuclear-density objects consisting of bound quarks or antiquarks over a significant range of masses. Here we analyze one such proposed scenario, which hypothesizes antiquark nuggets with a range of $B \\sim 10^{24-30}$ with specific predictions for spectral emissivity via interactions with normal matter. We find that, if these objects make up the majority of the dark matter density in the solar neighborhood, their radiation efficiency in solids is marginally constrained, due to limits from the total geothermal energy budget of the Earth. At allowed radiation efficiencies, the number density of such objects can be constrained to be well below dark matter densities by existing radio data over a mass range currently not restricted by other methods. ",
        "introduction": "Many different forms of astrophysical and cosmological evidence point to the existence of weakly interacting matter in an unknown form in the universe (see ~\\cite{DM0} for a recent review), with a total mass of order five times the inventory of what is observed as normal matter -- gas, stars, and dust -- giving a dark matter density estimated to be in the range $0.4-1$~GeV~cm$^{-3}$ in the solar neighborhood~\\cite{DM1,DM2,DM3} In addition, the current best models for  the early universe give stringent constraints on the content of interacting baryons in this dark matter at the time of nucleosynthesis~\\cite{DM0}.  The rather compelling arguments for non-baryonic dark matter have led to a wide variety of efforts, both theoretical and experimental, to either postulate or directly detect new particles, beyond the standard model, that would satisfy the dark matter characteristics. However, there remain several ``standard-model'' candidates for dark matter, which, if not currently favored, have not yet been excluded over all the possible range of parameters. In particular, very massive objects (at least compared to the $\\lesssim$TeV scale of typical particle candidates for dark matter) can still satisfy the astrophysical constraints on dark matter if their masses are sufficient that the flux in typical detectors is extremely low, but not so large that they are excluded due to galaxy dynamics or gravitational lensing observations. This may be translated into a constraint on their interaction cross sections per unit mass; current limits require $\\sigma/M \\lesssim 0.1$~cm$^2$~g$^{-1}$~\\cite{CSS08,Peter12}, a constraint that is in general easily satisfied by neutral objects with nuclear densities.  Such objects must not interact with normal matter via strong or electromagnetic channels at the time of nucleosynthesis. One candidate is the so-called quark nugget, or strangelet, hypothesized originally by Witten~\\cite{Witten84, Itoh70}, and developed in many variations since then. Quark nuggets can be neutral and metastable at their formation during the quantum chromodynamics phase transition of early-universe evolution, and thus do not undergo significant further interactions at nucleosynthesis, therefore evading the constraints on baryonic content.  Although still a matter for debate, the possibility of quark nugget {\\it color superconductivity}~\\cite{Wil98}, in which quarks near the Fermi surface of the nugget form correlated Cooper pairs, favors their possible stability~\\cite{LH04}. One of the more attractive aspects of these objects as candidates for dark matter is that the physics of their formation and interactions is in principle calculable according to the standard model, although such calculations can in practice be prohibitively difficult. ",
        "conclusions": ""
    },
    "1208/1208.4554_arXiv.txt": {
        "abstract": "We present the Transiting Exoearth Robust Reduction Algorithm (TERRA) --- a novel framework for identifying and removing instrumental noise in \\Kepler photometry.  We identify instrumental noise modes by finding common trends in a large ensemble of light curves drawn from the entire \\Kepler field of view.  Strategically, these noise modes can be optimized to reveal transits having a specified range of timescales.  For \\Kepler target stars of low photometric noise, TERRA produces ensemble-calibrated photometry having 33~ppm RMS scatter in 12-hour bins, rendering individual transits of earth-size planets around sun-like stars detectable as $\\sim 3 \\sigma$ signals. ",
        "introduction": "The \\Kepler Mission is ushering in a new era of exoplanet science.  Landmark discoveries include \\Kepler-10b, a rocky planet \\citep{Batalha:2011}; the \\Kepler-11 system of six transiting planets \\citep{Lissauer:2011el}; earth-sized \\Kepler-20e and 20f \\citep{Fressin12}; KOI-961b, c, and d -- all smaller than earth \\citep{Muirhead12}; and \\Kepler-16b a circumbinary planet \\citep{Doyle:2011ev}.  While \\Kepler has revealed exciting individual systems, the mission's legacy will be the first statistical sample of planets extending down to earth size and out to 1 AU.  \\Kepler is the first instrument capable of answering ``How common are earths?'' --- A question that dates to antiquity.  Planet candidates are detected by a sophisticated pipeline developed by the \\Kepler team Science Operations Center.  In brief, systematic effects in the photometry are suppressed by the Pre-search Data Conditioning (PDC) module, the output of which is fed into the Transiting Planet Search (TPS) module.  For further information, see \\cite{Jenkins:2010}.  The \\Kepler mission was designed to study astrophysical phenomena with a wide range of timescales, which include 1-hour transits of hot Jupiters, 10-hour transits of planets at 1 AU, and weeklong spot modulation patterns.  The PDC module is charged with removing instrumental noise while preserving signals with a vast range of timescales.  We review sources of instrumental errors in \\S~\\ref{sec:InstNoise}, highlighting the effects that are most relevant to transit detection.  The \\Kepler team has released candidate planets based on the first 4 and 16 months of data \\citep{Borucki11,Batalha12}.  Many of the candidates have additional followup observations from the ground and space aimed at ruling out false positive scenarios.  In addition, statistical arguments suggest that 90-95\\% of all candidates and that $\\sim 98$\\% of candidates in multi-candidate systems are bonafide planets \\citep{Morton:2011,Lissauer:2012}.  While \\textit{Kepler's} false positive rate is low, its completeness is largely uncharacterized.  If the completeness decreases substantially with smaller planet size or longer orbital periods, the interpretations regarding occurrence drawn from the \\cite{Borucki11} and \\cite{Batalha12} catalogs will be incorrect.  Hunting for the smallest planets, including earth-sized planets in the habitable zone, will require exquisite suppression of systematic effects.  Without optimal detrending, systematic noise will prevent the detection of the smallest planets, possibly the habitable-zone earth-sized planets, which is the main goal of the \\Kepler mission.  Therefore, it is essential for independent groups to develop pipelines that compliment both PDC and TPS.  An early example of an outside group successfully identifying new planet candidates is the Planet Hunters project \\citep{Fischer:2011bb,Lintott:2012ut}, which uses citizen scientists to visually inspect light curves.  In addition, existing pipelines from the HAT ground-based search \\citep{Huang:2012uj} and the \\textit{CoRoT} space mission \\citep{Ofir:2012va} have been brought to bear on the \\Kepler dataset yielding $\\sim100$ new planet candidates.  We present the Transiting Exoearth Robust Reduction Algorithm (TERRA) --- a framework for identifying and removing systematic noise.  We identify systematic noise terms by searching for photometric trends common to a large ensemble of stars.  Our implementation is tuned toward finding trends with transit-length timescales. ",
        "conclusions": "\\label{sec:conclusions} TERRA is a new technique for using ensemble photometry to self-calibrate instrumental systematics in \\Kepler light curves.  We construct a simple noise model by running a high-pass filter and removing thermal settling events before computing principle components.  For a typical $12.5 < $ \\Kmag $ < 13.5$ and CDPP12 $< 40$~ppm star, TERRA produces ensemble-calibrated photometry with 33~ppm RMS scatter in 12~hour bins.  With this noise level, a 100~ppm transit from an exoearth will be detected at $\\sim 3\\sigma$ per transit.  A potential drawback of removing thermal settling events is discarding photometry that contains a transit.  Thermal settling events amounted to 14\\% of the valid cadences in Q1-Q8 photometry.  Since signal to noise grows as the square root of the number of transits, removing 14\\% of the photometry results in a 7\\% reduction in the signal to noise of a given transit.  The completeness of the survey may decrease slightly, since some borderline transits will remain below threshold.  However, this can easily be overcome by gathering 14\\% more data.  Ensemble-based cotrending is most effective when the timescales in the ensemble are matched to the signal of interest. We are skeptical that a ``one size fits all'' approach exists and we encourage those who wish to get the most out of \\Kepler data to tune their cotrending to the timescale of their signals of interest."
    },
    "1208/1208.1824_arXiv.txt": {
        "abstract": "We present deep optical and X-ray follow-up observations of the bright unassociated $Fermi$-LAT gamma-ray source 1FGL J1311.7-3429. The source was already known as an unidentified EGRET source (3EG J1314-3431, EGR J1314-3417), hence its nature has remained uncertain for the past two decades. For the putative counterpart, we detected a quasi-sinusoidal optical modulation of $\\Delta m$ $\\sim$ 2 mag with a period of $\\simeq$1.5 hr in the $Rc$, $r'$ and $g'$ bands. Moreover, we found that the amplitude of the modulation and peak intensity changed by $\\gtrsim$ 1 mag and  $\\sim$0.5 mag respectively, over our total six nights of observations from 2012 March and May. Combined with $Swift$ UVOT data, the optical-UV spectrum is consistent with a blackbody temperature, $kT$ $\\simeq$ 1 eV, and the emission volume radius $R_{bb}$ $\\simeq$ 1.5$\\times$10$^4$ $d_{\\rm kpc}$ km ($d_{\\rm kpc}$ is the distance to the source in units of 1 kpc). In contrast, deep $Suzaku$ observations conducted in 2009 and 2011 revealed strong X-ray flares with a lightcurve characterized with a power spectrum density of $P(f)$ $\\propto$ $f^{-2.0\\pm0.4}$, but the folded X-ray light curves suggest an orbital modulation also in X-rays. Together with the non-detection of a radio counterpart, and significant curved spectrum and non-detection of variability in gamma-rays, the source may be the second ``radio-quiet'' gamma-ray emitting milli-second pulsar candidate after 1FGL J2339.7-0531, although the origin of flaring X-ray and optical variability remains an open question. ",
        "introduction": "The Large Area Telescope \\citep[LAT;][]{atw09} onboard the $Fermi$ Gamma-ray Space Telescope is a successor to EGRET onboard the Compton Gamma-ray Observatory \\citep{har99}, with much improved sensitivity, resolution, and energy range. The second \\F catalog, based on the first 24 months of all-sky survey data \\citep[2FGL;][]{2FGL}, provides source location, flux and spectral information, as well as light curves on month time bins for 1873 $\\gamma$-ray sources detected and characterized in the 100 MeV to 100 GeV range. Thanks to their small localization error circles (or ellipses) with typical 95$\\%$ confidence radii, $r_{95}$ $\\simeq 0.1\\deg-0.2\\deg$, for relatively bright sources, 69 $\\%$ of the 2FGL sources are reliably associated or firmly identified with counterparts of known or likely $\\gamma$-ray producing sources. In particular, more than 1,000 sources are proposed to be associated with AGN (of mainly the blazar class) and 87 sources with pulsars (PSRs), including 21 millisecond pulsars (MSPs) which are a new category of $\\gamma$-ray sources discovered with \\F \\citep{2FGL, MSP}. Other sources, albeit of a relative minority compared to AGNs and PSRs, also constitute important categories of new GeV sources like supernova remnants \\citep[SNRs;][]{W51, Cas-A, W44, IC443}, low-mass/high-mass binaries \\citep{LS5039, LSI, CygX-3}, pulsar wind nebula \\citep{Vela, 1509}, one nova \\citep{V407}, normal and starburst galaxies \\citep{LMC, M82}, and the giant lobes of a radio galaxy \\citep{cena-lobe}.   Despite such great advances in the identification of \\F sources, 575 (31$\\%$) sources in the 2FGL catalog still remain unassociated. Note that a substantial fraction of the unassociated sources (51$\\%$) have at least one analysis flag  due to various issues, while only 14$\\%$ of the associated sources have been flagged  \\citep{2FGL}. This may suggest some of unassociated sources are spurious due to complexity/difficulty of being situated in a crowded region near the Galactic plane. Nevertheless, many of them are bright enough to be listed in the one year \\F catalog \\citep[1FGL;][]{1FGL} and some of them are even listed in the bright source list based on the first 3 months of data \\citep[0FGL;][]{0FGL}. By comparing the distribution of associated and unassociated sources in the sky, a number of interesting features in the map were reported \\citep{2FGL}. For example, (1) the number of unassociated sources decreases with increasing Galactic latitude, (2) the number of unassociated sources increases sharply below Galactic latitudes, $|b|$ $<$ 10$^{\\circ}$, and (3) the fraction of sources with curved gamma-ray spectra among the unassociated sources is greater (28$\\%$) than the fraction of curved spectra sources among the associated sources (16$\\%$). Further extensive studies based on a statistical approach in an effort to correlate their gamma-ray properties with the AGN and  PSR populations was presented for 1FGL unassociated sources \\citep{UnIDstat}.   In this context, 1FGL J1311.7-3429 (or 2FGL name, 2FGL J1311.7-3429) is a $classical$ unassociated gamma-ray source situated at high Galactic latitude ($l$ = 307.6859\\deg, $b$ = 28.1951\\deg), and was first discovered by EGRET about 20 years ago as 3EG J1314-3431 \\citep{har99} or EGR J1314-3417 \\citep{cas08}. The source was also reported by \\F in the 0FGL list with a gamma-ray flux $F_{\\rm 0.1-20 GeV}$ = (11.7 $\\pm$1.1) $\\times$ 10$^{-8}$ photons cm$^{-2}$ s$^{-1}$, which is marginally consistent with the gamma-ray flux determined by EGRET, $F_{\\rm 0.1-20 GeV}$ = (18.7 $\\pm$3.1) $\\times$ 10$^{-8}$ photons cm$^{-2}$ s$^{-1}$, within the 2$\\sigma$ level. In the 2FGL catalog, the detected significance of 1FGL J1311.7-3429 is 43.1 $\\sigma$, which is one of the brightest sources with an \\textsc{unid} flag.  Based on the one-month binned $>100$ MeV gamma-ray light curve of 1FGL J1311.7-3429 over two years of data as published in the 2FGL catalog\\footnote{http://heasarc.gsfc.nasa.gov/FTP/fermi/data/lat/catalogs/source/lightcurves/\\\\ 2FGL$\\_$J1311d7m3429$\\_$lc.png}, no statistically significant variability was observed (\\textsc{Variability\\_index} = 19.09; \\cite{2FGL}). The gamma-ray spectrum is significantly curved with curved significance \\textsc{Signif\\_curve} of 6.33 \\citep{2FGL}. \\footnote{As detailed in \\cite{2FGL}, the \\textsc{Variability\\_index} is an indicator of the flux being constant across the full 2-year period. A value of \\textsc{Variability\\_index} $>$ 41.6 is used to identify variable sources at a 99$\\%$ confidence level. Similarly, the \\textsc{Signif\\_curve} parameter is an indicator of the spectrum being curved, by comparing the likelihood values calculated for a LogParabola and a single power-law function. \\textsc{Signif\\_curve} is distributed as $\\chi^2$ with one degree of freedom, and \\textsc{Signif\\_curve} $>$ 16 corresponds to 4 $\\sigma$ significance of curvature.}  The first X-ray follow-up observation of 1FGL J1311.7-3429 was conducted as a part of $Suzaku$ X-ray observations of 11 unidentified \\F objects at high Galactic latitude, $| b |$ $>$ 10$^{\\circ}$ \\citep{mae11,tak12}. The X-ray source associated with 1FGL J1311.7-3429 showed a very rapid X-ray flare with the count rate changing by a factor of 10. Subsequent $Chandra$ ACIS-I (2010 March 21 for a 19.87 ks  exposure, obsID 11790) and $Swift$ XRT observations (2009 February 27 for a 3.34 ks exposure, obsID 31358; see also Table 1) confirmed that the brightest X-ray source within the \\F error ellipse is the most credible counterpart and that the X-ray source is also variable on month-to-year timescales. The unabsorbed X-ray flux observed with $Chandra$ was 1.03$\\times$ 10$^{-13}$ erg cm$^{-2}$ s$^{-1}$ in the 0.5$-$8 keV band, with a differential photon spectral index, $\\Gamma$ = 1.26$^{+0.38}_{-0.37}$ \\citep{che12}.  Motivated by the initial X-ray results, we conducted further deep observations of 1FGL J1311.7-3429 with  $Suzaku$, together with deep optical observations using a 105 cm Ritchey-Chr\\'etien telescope ($g'$, $Rc$, and $Ic$ bands) at the Ishigakijima Astronomical Observatory (IAO) in Japan, as well as the 1m telescope ($r'$ band) at Lulin Observatory in Taiwan. In Section 2, we describe the details of the $Suzaku$ observations and optical observations and data reduction procedures. Very recently, \\cite{rom12} reported quasi-sinusoidal optical modulation of this source with a 1.56 hr (5626 s) period, suggesting that the source is another black-widow-type MSP like that recently discovered for 1FGL J2339.7-0531 \\citep{rom11,kon12}. Our paper confirms some of those optical findings for 1FGL J1311.7-3429, plus provides results from multiple epoch optical monitoring between 2012 March and May and completely new X-ray data based on a long $Suzaku$ observation conducted in 2011 together with our previously published archival 2009 data. The results of these observations are given in Section 3. Based on our new observational data in optical and X-ray, and various observed gamma-ray parameters compiled in the 2FGL source catalog, we support that 1FGL J1311.7-3429 could be a ``radio-quiet'' gamma-ray emitting MSP candidate like 1FGL J2339.7-0531. The variable optical/X-ray source is posited as the counterpart to the gamma-ray source and throughout, we refer to it simply as 1FGL J1311.7-3429.  \\begin{center} \\begin{deluxetable*}{lcccc}{ht} \\tablecaption{Observation log of 1FGL J1311.7-3429 analyzed in this paper} \\tablewidth{0pt} \\tablehead{ \\colhead{Obs Start (UT)} &  \\colhead{Obs End (UT)} & \\colhead{Observatory} & \\colhead{Band} & \\colhead{Exposure} } \\startdata 2009-08-04 04:56:35 & 2009-08-05:07:18:14  & $Suzaku$ XIS & X-ray & 33.0 ks \\\\ 2011-08-01 16:48:20 & 2011-08-03:17:40:15  & $Suzaku$ XIS & X-ray & 65.2 ks \\\\ 2009-02-27 18:41:59 & 2009-02-27:22:04:57  & $Swift$ UVOT & UV & 276/276/276/552/742/1106$^a$ \\\\ 2012-03-24 17:45:37 & 2012-03-24 19:38:35  & LOT 1m  & $r'$ & 5 min $\\times$ 21 \\\\ 2012-03-25 17:48:26 & 2012-03-25 19:26:33  & LOT 1m  & $r'$ & 5 min $\\times$ 12 \\\\ 2012-03-26 17:25:14 & 2012-03-26 18:29:59 & LOT 1m  & $r'$ & 5 min $\\times$ 12 \\\\ 2012-04-12 16:35:08 & 2012-04-12 18:24:30 & LOT 1m  & $r'$ & 5 min $\\times$ 21 \\\\ 2012-05-24 13:00:49 & 2012-05-24 15:34:44  & LOT 1m  & $r'$ & 5 min $\\times$ 29 \\\\ 2012-05-25 12:31:11 & 2012-05-25 15:41:11 & IAO 105cm & $g'$, $Rc$, $Ic$ & 20 min $\\times$ 7  \\\\  \\tableline \\enddata \\tablecomments{$^a$: $Swift$ UVOT exposures for the $v/b/u/uvw1/uvm2/uvw2$ bands in seconds.} \\end{deluxetable*} \\end{center} ",
        "conclusions": "During the first (AO4; 2009) and second (AO6; 2011) $Suzaku$ observations, we detected significant variability in 1FGL J1311.7-3429 characterized by repeated flaring activity, with a timescale of $\\sim$ 10 ks. The variability is only clearly seen above 1 keV. The NPSD is well characterized by $P(f)$ $\\propto$ $f^{-2}$, as is the case for the X-ray variability of various classes of AGN including Seyferts and blazars \\citep[e.g.,][]{hay98, vau03}. The latter class constitutes the majority of \\F sources but the X-ray variability timescales of blazars are in general somewhat longer,  typically $\\simeq$ 1 day \\citep[e.g.,][]{kat01}. Such ``red-noise'' PSD behavior is also observed in X-ray (Galactic black hole and neutron star) binary systems, but on much shorter time scales. For example, variability as short as $\\sim$ 1$-$10 ms has been observed for the famous Galactic black hole source Cyg X-1 \\citep[e.g.,][]{mee84, hay98}.   \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=0.47]{Fig9a.ps} \\includegraphics[angle=0,scale=0.47]{Fig9b.ps} \\caption{($top$) Broadband spectrum of 1FGL J1311.703429. The X-ray data are as presented in Figure 8. The gamma-ray data points are taken from the 2FGL catalog \\citep{2FGL}. The radio upper limit of 2.5 mJy at 1.4 GHz is taken from the NVSS catalog \\citep{NVSS}. The optical and UV data represents $Swift$ UVOT data (see Maeda et al. 2011 and Cheung et al. 2012). ($bottom$) Close-up of the broadband spectrum in the optical-UV range. Epochs 0$-$6 (1200 s each) correspond to time ranges defined in Figure 2. Note the optical-UV spectrum is well fitted by a blackbody model of $kT$ $\\sim$ 1eV with a radius of the emission  volume of $R_{bb}$ $\\simeq$ 1.5$\\times$10$^4$ $d_{\\rm kpc}$ km, where $d_{\\rm kpc}$ is distance to the source in unit of 1 kpc.} \\end{center} \\end{figure}   Significant variability has also been observed in the optical ($g'$, $Rc$, $r'$), where it is rather a quasi-sinusoidal flux modulation with a 1.56 hr period, as recently reported by \\citet{rom12}. Moreover, we also found that the modulation profile, including the amplitude of modulation and peak intensity, has changed largely among the six nights of observations. The apparent modulation of magnitude observed in both the IAO and LOT data is quite similar to those observed in 1FGL J2339.7-0531, which is characterized by a 4.63-hr orbital period in optical and X-ray data \\citep{rom11,kon12}. Note that 1FGL J2339.7-0531 is now suggested to be a ``radio-quiet'' gamma-ray emitting black-widow MSP with a $\\simeq$0.1 $M_{\\odot}$ late-type companion star, viewed at inclination $i$ $\\simeq$ 57$^{\\circ}$. Moreover, this compact object companion in 1FGL J2339.7-0531 is strongly heated, with $T_{\\rm eff}$ varying from $\\sim$ 6900 K (superior conjunction) to $<$ 3000 K  at minimum \\citep{rom11}.  \\begin{figure} \\begin{center} \\includegraphics[angle=0,scale=0.53]{Fig10a.ps} \\includegraphics[angle=0,scale=0.53]{Fig10b.ps} \\caption{($top$) Comparison of the 2FGL Variability Index versus Curvature Significance for associated sources (AGN: $blue$, PSR: $green$) and unassociated sources ($red$). ($bottom$) Comparison of the 2FGL Photon Index versus $0.1-100$ GeV Flux for the associated sources and unassociated sources. In both panels, a separation between the AGN and PSR populations is evident. Note that 1FGL J1311.7-3429 is situated in the typical PSR region in the $top$ panel, whilst at the boundary of AGN/PSR sources in the $bottom$ panel. } \\end{center} \\end{figure}    In this context, the spectral energy distribution of 1FGL J1311.7-3419 from radio to gamma-rays (Figure 9, $top$) may provide some hints on the nature of this mysterious source.  1FGL J1311.7-3419 has been the subject of both radio pulsar counterpart searches and blind searches for gamma-ray pulsations, but to date no pulsed emission has been detected \\citep{ran11}, and the situation is quite similar to the case for 1FGL J2339.7-0531. There are no NVSS radio sources catalogued within the 2FGL ellipse of 1FGL J1311.7-3429 down to 1.4 GHz flux completeness limit of $\\simeq$2.5 mJy \\citep{NVSS}, which is given as an arrow in Figure 9 ($top$). From the 0.1$-$100 GeV gamma-ray flux listed in the 2FGL catalog, $F_{\\gamma}$ $\\equiv$ $F_{\\rm 0.1-100GeV}$ $\\simeq$ 6.2$\\times$10$^{-11}$ erg cm$^{-2}$ s$^{-1}$,  we obtain $F_{\\gamma}$/$F_X$ $\\simeq$ 300, $F_{\\gamma}$/$F_R$ $\\ge$ 1.7$\\times$10$^6$ for 1FGL J1311.7-3429, where $F_X$ and $F_R$ are the X-ray and radio fluxes measured in 2$-$10 keV and at 1.4 GHz, respectively. Note, these values are almost comparable to those measured in 1FGL J2339.7-0531,  with $F_\\gamma$ $\\simeq$ 3.0$\\times$10$^{-11}$ erg cm$^2$ s$^{-1}$, $F_{\\gamma}$/$F_X$ $\\simeq$ 150 and $F_{\\gamma}$/$F_R$ $\\ge$ 8.7$\\times$10$^5$. Although its flat spectrum X-ray continua and relatively high $\\gamma$-ray-to-X-ray energy flux ratio (of $\\gtrsim$ 100) is typical of the sub-class of blazars \\citep[FSRQ;][]{2FGL}, the non-detection in the radio as well as its curved gamma-ray spectrum seems to strongly disfavor an AGN association for 1FGL J1311.7-3429.  To further support the MSP association of 1FGL J1311.7-3429, Figure 9 ($bottom$) shows a close-up of the SED of 1FGL J1311.7-3429 in the optical/UV bands, reconstructed from IAO and $Swift$ UVOT data. Observed magnitudes were converted to flux density based on \\cite{fuk95}. $Red$ $dashed$ $line$ indicates a tentative fit with a blackbody model of $T$ $\\simeq$ 1.2$\\times$10$^4$ K (or $kT$ $\\simeq$ 1 eV) assuming an emission volume radius of $R_{bb}$ $\\simeq$ 1.5$\\times$10$^4$ $d_{\\rm kpc}$ km, where $d_{\\rm kpc}$ is the distance to the source in units of 1 kpc. Therefore, the optical/UV spectrum of 1FGL J1311.7-3429 seems compatible with what is expected from a companion star of a radio-quiet MSP like 1FGL J2339.7-0531. A slightly higher temperature than 1FGL J2339.7-0531 \\citep[$kT$ $\\simeq$ 0.3$-$0.6 keV;][]{rom11} may indicate $\\times$2 smaller orbital radius, $\\sim$ 5$\\times$10$^5$ km, for the 1FGL J1311.7-3429 binary system, assuming that the mass of the companion star is 0.1 $M_{\\odot}$ and the pulsar spin-down luminosity is  $L$ $\\simeq$ 10$^{34}$ erg s$^{-1}$ \\citep[parameters suggested for 1FGL J2339.7-0531 binary system;][]{rom11}. A change of the modulation profile observed in the $r'$-band may be accounted for by rapid changes in the companion star temperature, but this remains uncertain.  Since both the optical and X-ray light curves in the case of 1FGL J2339.7-0531 clearly exhibits a 4.63-hr orbital modulation \\citep{rom11,kon12}, the detection of periodicity in the X-ray light curve of 1FGL J1311.7-3429 is also likely. In fact, the folded X-ray light curve exhibits an excess feature around $\\phi$ = 0.5. As indicated by \\cite{rom12}, this is presumably pulsar superior conjunction, although both the normalized intensity as well as the phase peak appear to have changed substantially between the 2009 (AO4) and 2011 (AO6) observations. We therefore expect that X-ray variability consists of at least two different components -- one associated with the binary motion as for the optical data, and the other is rather random fluctuation well represented by the NPSD of $P(f)$ $\\propto$ $f^{-2}$, whose physical origin is still unknown but possibly related with perturbations associated with shock acceleration.  Such flaring X-ray variability has not yet been observed for 1FGL J2339.7-3429, but solely in 1FGL J1311.7-3429. In a review of X-ray emission from MSPs, \\cite{zav07} describes three primary sources of X-ray emission: (1) intra-binary shock, (2) the neutron star (NS) itself, and (3) pulsar wind nebula outside the binary system \\citep[see, also][]{arc10}. In fact, thermal emission of $kT$ $\\sim$0.1$-$0.2 keV is often observed from MSPs, which is thought to arise from the surface of the NS and is steady with time \\citep[e.g.,][]{mar11,mae11}. More recently, thermal emission of $kT$ $\\simeq$ 0.1 keV was also detected from 1FGL J2339.7-0531 (Kong et al. 2012, in prep), which is again steady despite large flux modulation associated with binary motion being observed above 2 keV.  Therefore, the fact that X-ray variability is not clearly seen below 1 keV may suggest there could be some contribution from the surface of the NS, although from the spectral fitting, it is not statistically significant. Then the variable, hard X-ray emission could arise from the intra-binary shock rather than the nebula because variability as short as $\\simeq$ 10 ks is unlikely to originate from the extended pulsar wind nebula\\footnote{Variability has been observed in some pulsar wind nebula in X-rays and gamma-rays, but with typical timescales of a week to months \\citep[e.g.,][]{pav01,Crab}.}. Such a shock  could readily produce gamma-ray emission \\citep{aro93}. If localized, it could easily account for the orbital modulation as seen in the X-ray emission from 1FGL J2339.7-0531. But if material leaving the companion star, either through Roche-lobe overflow or a stellar wind is non-uniform or patchy, we might expect random flaring activity as we see in the X-ray data of 1FGL J1311.7-3429.  We can also speculate on the nature of 1FGL J1311.7-3429 based solely on the gamma-ray properties, an approach already applied to the 1FGL unIDs in \\cite{UnIDstat}. Figure 10 ($top$) presents a comparison of the 2FGL associated (either AGN ($blue$) or PSRs ($green$)) and unassociated sources ($red$) in the \\textsc{Variability\\_index} and \\textsc{Signif\\_curve} plane. Apparently, 1FGL J1311.7-3429 is situated in the typical PSR regions of this diagnostic plane. Similarly, Figure 10 ($bottom$) plots the distribution of PSRs and AGN in the \\textsc{Photon index} versus \\textsc{F$_{\\rm 0.1-100 GeV}$} plane. In this case, 1FGL J1311.7-3429 is at the boundary of AGN and PSR sources, so is still consistent with a PSR association. For comparison, we also plot the gamma-ray parameters for 1FGL J2339.7-0531. In both panels, the gamma-ray properties of 1FGL J1311.7-3429 and 1FGL J2339.7-0531 are quite similar. Again, this supports the idea that 1FGL J1311.7-3429 is a black-widow system and may be a second example of a ``radio-quiet'' MSP after 1FGL J2339.7-0531.  Finally, if the rapid X-ray flaring variability observed with $Suzaku$ may be due to inhomogeneity of shock material and/or rapid changes in the beaming factor, this could be expected to occur also in gamma-rays, as suggested by a smooth connection of the spectrum between X-ray and gamma-ray energies. Moreover, we may see correlated variability also in the optical, which may be related to the change of modulation profile in the  $r'$-band magnitudes we see in Figure 3. Unfortunately, such fast variability is difficult to observe in gamma-rays as we referred to one-month binned light curve in Section 1, despite the excellent sensitivity of $Fermi$-LAT. The low gamma-ray statistics also make it difficult to run a cross-correlation in order to measure any possible correlated optical, X-ray and gamma-ray variability. Further continuous investigation is necessary to confirm the origin of ``variable'' X-ray emission observed in the 1FGL J1311.7-3429 system."
    },
    "1208/1208.2685_arXiv.txt": {
        "abstract": "We propose an economical model to explain the apparent 130 GeV gamma ray peak, found in the Fermi/LAT data, in terms of dark matter (DM) annihilation through a dipole moment interaction. The annihilating dark matter particles represent a subdominant component, with mass density $7-17$\\% of the total DM density; and they only annihilate into $\\gamma\\gamma$, $\\gamma Z$, and $ZZ$, through a magnetic (or electric) dipole moment. Annihilation into other standard model particles is suppressed, due to a DM mass splitting in the magnetic dipole case, or to $p$-wave scattering in the electric dipole case. In either case, the observed signal requires a dipole moment of strength  $\\mu\\sim 2/$TeV. We argue that composite models are the preferred means of generating such a large dipole moment, and that the magnetic case is more natural than the electric one.  We present a simple model involving a scalar and fermionic techniquark of a confining SU(2) gauge symmetry.  We point out some generic challenges for getting such a model to work. The new physics leading to a sufficiently large dipole moment is below the TeV scale, indicating that the magnetic moment is not a valid effective operator for LHC physics, and that production of the strongly interacting constituents, followed by techni-hadronization, is a more likely signature than monophoton events. In particular, 4-photon events from the decays of bound state pairs are predicted. ",
        "introduction": "\\label{coann} The annihilation cross section for $\\chi_1\\chi_2\\to f\\bar f$ for a standard model fermion pair at low velocity, via a magnetic moment interaction, can be expressed as \\be \\sigma_{\\bar f f} v = \\alpha\\mu^2 \\sum_i N_i \\ee where $N_i=1$ for a hypothetical particle with unit electric charge and no coupling to the $Z$.   The sum is over all kinematically accessible final states.   For standard model fermions, and similarly parametrizing the contributions from $W^+W^-$ and $Zh$ final states, $N_{\\rm eff}$ is \\be N_i = \\left\\{ \\begin{array}{ll} q_i^2(1-2v_it_\\W\\xi +(v_i^2+a_i^2)t_\\W^2\\xi^2)), & f \\bar f\\\\ {1\\over 16}\\left(m_Z\\over m_W\\right)^4 \\xi^2\\,\\psi_{\\W\\W},& WW\\\\ {1\\over 16\\, c_\\W^4}\\,\\xi^2\\,\\psi_{hZ},& hZ \\end{array}\\right. \\ee The $\\bar f f$ result is as given in ref.\\ \\cite{Weiner:2012cb}, where $q_i$ is the electric charge, $v_i,a_i$ are the vector and axial-vector couplings respectively to the $Z$, in units of $q_i e$, and $\\xi = (1-m^2_Z/4(m^2_\\chi))^{-1}$.  For $WW$ we define $\\psi_{\\W\\W} = \\left(1 + 4\\epsilon_\\W - {17\\over 4} \\epsilon_\\W^2 -\t\\frac34\\epsilon_\\W^3\\right)(1-\\epsilon_\\W)^{1/2}$ with $\\epsilon_\\W = m_W^2/m_\\chi^2$. We find that the contribution from neutrinos is $\\sum_{i} N_i = 3\\xi^2/(8c_\\W^4)$, and for all charged leptons or quarks $v_i=t_\\W((4s_\\W^2|q_i|)^{-1}-1)$, $a_i = t_\\W/(4s_\\W^2|q_i|)$, giving $\\sum_i N_i = 10.4$ from fermionic final states, assuming $m_\\chi = 130$ GeV and $s_\\W^2 = 0.23$.  Because of a strong cancellation between the virtual $\\gamma$ and $Z$ contributions to the $WW$ channel, it contributes only $0.20$ to $\\sum_i N_i$. For $hZ$, we define $\\psi_{hZ} = \\left(1-\\frac12(\\epsilon_h - 5\\epsilon_\\Z) +\\frac{1}{16}(\\epsilon_h-\\epsilon_\\Z)^2\\right)\\times$ $\\left(1-\\frac12(\\epsilon_h+\\epsilon_\\Z) + \\frac{1}{16}(\\epsilon_h-\\epsilon_\\Z)^2\\right)^{1/2}$ with $\\epsilon_h = m_h^2/m_\\chi^2$, contributing $0.22$ to $\\sum_i N_i$. Hence the total is $\\sum_i N_i = 10.8$.  For the $f\\bar f$ channels, the result from an electric dipole moment interaction is the same, except for an additional velocity-squared suppression factor of $v^2/3$. ",
        "conclusions": ""
    },
    "1208/1208.0155_arXiv.txt": {
        "abstract": "{The binary B5\\,V star BD$+$31$\\degr$643 exhibits a disk-like structure detected at optical wavelengths. Even though the feature is well centered on the star, it has been argued, based on \\emph{Spitzer } observations, that the feature is a filament not directly associated to the binary star.} {The purpose of the present paper is to investigate whether polarization imaging may provide evidence either for or against the disk hypothesis. In addition,  we aim at clarifying whether there might be any additional close companion to the binary star.} {We used the coronagraph \\emph{PolCor} in its polarization mode in combination with an EMCCD camera allowing short unit exposure times. As a result of shift-and-add and frame selection, the spatial resolution is improved compared to traditional CCD imaging. In order to possibly reveal an additional stellar companion, we used high resolution spectroscopy in the optical and high spatial resolution imaging in the near-IR. } {The disk/filament is much better seen in polarization; it is narrow and a line drawn along the ridge passes within a second of arc from the star. The degree of polarization is high ($\\approx$50\\% after correction for the extended component of the reflection nebula) which means that the disk/filament must be approximately at the same distance as the star. Although we confirm that the feature is much brighter south-east than north-west of the star, the evidence that the feature is physically connected to the star is strengthened and suggests that we are witnessing the destruction process of an accretion disk. Our spectroscopy shows that at least one of the stars is a spectroscopic binary. We were, however, not able to spatially resolve  any stellar component in addition to the two well separated stars.   } {} ",
        "introduction": "IC 348 is a young star cluster, located next to the Perseus molecular cloud complex, with its most luminous member -- the B5$\\,$V binary BD$+$31$\\degr$643 --  partially embedded to create a prominent reflection nebula visible both in the optical and near-infrared (NIR). Although the cluster did not necessarily form in the cloud, its most massive star(s) is undoubtedly interacting with it through stellar winds and UV radiation \\citep[][ and references therein]{Arce2010}. \\emph{Spitzer} maps reveal that the IR emission surface brightness peak of the cloud coincides with the location of BD$+$31$\\degr$643 \\citep{Schmeja2008,Muench2007}, and that an incomplete spherical shell (with a diameter of $\\sim$200{\\arcsec}) surrounds the binary \\citep{Rebull2007}.  One of the most debated features in the nebulosity is the bright linear streak first imaged with optical coronagraphy at the University of Hawaii 2.2\\,m telescope by \\citet{Kalas1997}, centred on BD$+$31$\\degr$643, with an extent of 6$\\,$600$\\,$AU. It was originally believed to originate from light scattered in a huge circumstellar debris disk, but subsequent imaging of the region has cast doubt on the nature of the streak. Although a narrow disk-like feature with the same position angle across the star was seen with \\emph{Spitzer} in extended 24{$\\,\\upmu$}m emission, \\citet{Rebull2007} dismissed the disk theory based on several arguments: 1) the scattered optical disk would be an order of magnitude larger than any known debris disk; 2) the size at 24$\\,\\upmu$m is three times bigger that the optical size; 3) no excess IR \\citep[or submillimetre,][]{Enoch2006} emission is seen in the spectral energy distribution (SED) of the star \\citep[although initially suggested by ISO SWS observations by][the spectrometer's large aperture was probably contaminated by extended emission]{Merin2004b}; 4) extreme 35{\\degr} opening angle toward the south-east (SE); and 5) strong brightness asymmetry. The 24$\\,\\upmu$m surface brightness peaks 22{\\arcsec} SE of the star and is four times brighter than on the corresponding north-west (NW) position. In addition, it extends to 55{\\arcsec} on the SE side, but reaches 80{\\arcsec} on the NW side.  \\citet{Rebull2007} instead suggest that the linear structure observed in the cavity cleared by BD$+$31$\\degr$643 may be caused by a filamentary interstellar cloud---a common substructure in star forming molecular clouds \\citep[see, e.g.,][]{Williams2000} -- which coincidentally crosses behind or in front of the star. The brightness asymmetry would be explained by the filament's viewing angle and closer approach to the SE side of the star.  The physical properties of the star BD$+$31$\\degr$643 are summarized in Table\\,\\ref{table:data}.   \\begin{table} \\caption{Coordinates and physical properties of BD$+$31$\\degr$643 (also commonly denoted HD 281159).} \\label{table:data} \\centering \\begin{tabular}{ l l l}     % \\hline\\hline Parameter & Value & Ref.\\\\ \\hline Position (J2000) & R.A.= 03h$\\,$44m$\\,$34.19s & 1\\\\ & Dec. = +32{\\degr}09{\\arcmin}46.14{\\arcsec} & 1 \\\\ Distance\\tablefootmark{a}, $D$ & $300\\,$pc & 2 \\\\ Spectral type and luminosity class & B5$\\,$V \\& B5$\\,$V binary &  \\\\ Effective temperature, $T_{\\textnormal{eff}}$ & $15\\,400\\,$K & 3 \\\\ Stellar mass, $M_{*}$ & $6.0{\\pm}1.0\\,M_{\\odot}$ & 4 \\\\ Surface gravity, log$\\,g$ & $3.75{\\pm}0.05\\,$cm$\\,$s$^{-2}$ & 5 \\\\ Age & 1.3$\\,$Myr & 4 \\\\ \\hline \\end{tabular} \\tablebib{ (1)~vanLeeuwen2007; (2)~\\cite{Herbst2008}; (3)~\\citet{Merin2004b}; (4)~\\citet{Preibisch2001}; (5)~\\citet{Montesinos2009} } \\tablefoot{ \\tablefoottext{a}{The distance measured from the mean parallax of IC$\\,$348 cluster members is  ($\\sim$260$\\,$pc), while the distance  derived from the HR diagram is  greater ($\\sim$320$\\,$pc) \\citep[see][and references therein]{Haisch2001}.}} \\end{table}  \\subsection{Aims of this study} The scattered light imaged by \\citet{Kalas1997} and the extended IR emission  seen by \\citet{Rebull2007} both reveal a disk-like feature crossing the binary star BD$+$31$\\degr$643, but results  remain inconclusive on whether or not it actually is a disk. Multicolor \\emph{aperture} polarimetry has shown a polarization that is parallel to the major axis of the projected disk feature and perpendicular on the sky to the projected direction of the local magnetic field \\citep{Andersson1997}. Our goal was to make the first polarimetric \\emph{mapping} of the structure, with the highest spatial resolution and contrast to date, using polarimetric coronagraphy with PolCor at the 2.6-m Nordic Optical Telescope (NOT), complemented with high-resolution optical and IR spectra, and high spatial resolution IR imaging, in order to discriminate between its filamentary or circumstellar disk nature, and possibly resolve additional stellar components. ",
        "conclusions": "\\label{sec:con}  We have shown that the disk-like feature exhibits a large contrast in polarized light and the high polarization degree, achieved after subtraction of the background, shows that the dust giving rise to the feature must be located in the vicinity of the star. A line drawn along the ridge of the feature passes very close to the direction of the star, which in principle can be a coincidence but taken together, these two findings indicate a physical connection between the feature and the central binary star. However, if the feature is due to a circumstellar ring it must have a non-uniform density distribution with a sector NW of the stars having a void of dust. If, after all, the feature is due to a narrow filament we show that the polarization results can be explained assuming a certain orientation of the filament.  The bright SE part of the disk-like feature as seen in polarized light perfectly coincides with the ISO/CVF image at 16\\,$\\mu$m.  The feature is, however not seen above the bright background in any of the PAH peaks at shorter wavelengths. This fact indicates that any PAH molecules that may have been hosted in the dust feature have disappeared, probably blown away by the radiation pressure. This means that the feature must have been close to the star long enough for such a process to be efficient. We have not estimated the required time scale, but it may add an additional constraint favoring the ring hypothesis.  We show that at least one of the visual binary stars is a spectroscopic binary and that each of the SB components have $v\\sin(i)$ values typical for B stars. The velocity separation at the time of our high resolution spectroscopy was large, 307 km/s, indicating that the SB should have a short period (days) and our intermediate spectroscopy during five consecutive nights support this interpretation. The system should be monitored spectroscopically (slit spectroscopy, including the two binary stars) for assessing the orbital parameters."
    },
    "1208/1208.1129_arXiv.txt": {
        "abstract": "The super-massive 4 million solar mass black hole (SMBH) SgrA* shows variable emission from the millimeter to the X-ray domain. A detailed analysis of the infrared light curves allows us to address the accretion phenomenon in a statistical way. The analysis shows that the near-infrared flux density excursions are dominated by a single state power law, with the low states of SgrA* limited by confusion through the unresolved stellar background. We show that for 8-10m class telescopes blending effects along the line of sight will result in artificial compact star-like objects of 0.5-1 mJy that last for about 3-4 years. We discuss how the imaging capabilities of GRAVITY at the VLTI, LINC-NIRVANA at the LBT and METIS at the E-ELT will contribute to the investigation of the low variability states of SgrA*. ",
        "introduction": "\\label{sec:intro}  %  The thorough investigation of stellar number densities, luminosities and orbits as well as light curves of SgrA* allows a detailed analysis of the immediate surroundings of the super massive black hole at the center of the Milky Way. However, these investigations also have shown that there are limits for the currently employed methods. Determining the K-band luminosity function (KLF) of the S-star cluster members from infrared imaging and using  the shape of the stellar distribution from number density counts, Sabha et al. (2012) have been able to shed some light on the amount and nature of the stellar and dark mass that is or may be associated with the cluster of high velocity S-star cluster in the immediate vicinity of Sgr~A*. However, for both the identification of faint stars and the investigation of faint flux density levels of SgrA* the stellar confusion limit as present for 8-10m class telescopes imposes some of these limits.  In addition to problems simply arising from the stellar confusion limit, there are other limitations that arise from aliasing populations of massive objects or disturbing emission mechanisms or confusing emission from particular source structures. For completeness we mention these effects here as well. There will for instance be source components in the extended accretion flow and the thermal Bremsstrahlung component (SgrA* is located in) that add confusion to the determination of the intrinsic variability of SgrA* in the radio and X-ray domain, respectively.  It is also expected that there is a population of dark stellar remnants at the bottom of the potential well. This is based on dynamical and stellar evolutionary arguments (see e.g. Morris et al. 1993) with some observational evidence through the increase of the number density of X-ray binaries many of which harbor a stellar black hole as a companion (Muno et al. 2005). Using Monte Carlo simulations of two-body relaxation, tidal disruptions of stars by Sgr~A* in addition to direct plunges through the event horizon Freitag et al. (2006) find that within 0.01, 0.1 and 1~pc from the Center there are approximately 560, $2.4 \\times 10^4$ and $2.1 \\times 10^5$~\\solm in 10~\\solm stellar black holes and 180, 6500 and $3.4 \\times 10^5$~\\solm in main-sequence stars. In addition roughly the same amount of white dwarfs and neutron stars are expected. The interaction between the more massive of these stellar remnants and the visible high velocities of the central S-star cluster will have an influence on the orbital parameters of these sources.  In this article we describe and summarize some aspects of these effects and point at the requirements and properties of upcoming instrumentation that will, at least partially, help to overcome some of these confusion problems. Here we concentrate on the role of LINC-NIRVANA (Herbst et al. 2010, Horrobin et al. 2010) at the LBT (Vaitheeswaran et al. 2010), GRAVITY at the VLTI (Eisenhauer et al. 2011, Straubmeier et al. 2010), and METIS (Brandl et al. 2010) at the E-ELT (Gilmozzi, R. \\& Spyromilio, J., 2008). For GRAVITY at the VLTI the positioning (phase-referencing) accuracy will be of the order of 10~$\\mu$as at 2$\\mu$m wavelength (infrared K$_s$-band). The GRAVITY imaging resolution in this band will be $\\le$4~mas, for LINC-NIRVANA at the LBT the angular resolution will be better than 20~mas in the same band, and  at 3.8~$\\mu$m for METIS at the E-ELT the angular resolution will be about 20~mas as well. ",
        "conclusions": ""
    },
    "1208/1208.0858_arXiv.txt": {
        "abstract": "There are insufficient super soft ($\\sim 0.1$ keV) X-ray sources in either spiral or elliptical galaxies to account for the rate of explosion of Type Ia supernovae in either the single degenerate or the double degenerate scenarios. We quantify the amount of circumstellar matter that would be required to suppress the soft X-ray flux by yielding a column density in excess of $10^{23}$ cm$^{-2}$. We summarize evidence that appropriate quantities of matter are extant in SN~Ia and in recurrent novae that may be supernova precursors. The obscuring matter is likely to have a large, but not complete, covering factor and to be substantially non-spherically symmetric. Assuming that much of the absorbed X-ray flux is re-radiated as black-body radiation in the UV, we estimate that $\\lesssim$ 100 sources might be detectable in the \\galex\\ all sky survey. ",
        "introduction": "The first suggestion that Type Ia (SN~Ia) supernovae may arise through mass transfer from a non-degenerate star onto a white dwarf may have been Wheeler \\& Hansen (1971). This suggestion of a single degenerate (SD) model was quantified by Whelan \\& Iben (1975) and pursued by many since. An alternative model is the merger of two degenerate stars, the double-degenerate (DD) model (Iben \\& Tutukov 1984; Webbink 1984). Both of these models are constrained by the paucity of bright, soft, X-ray sources (Di Stefano 2010a,b). In the SD model, an associated constraint is the strong expectation that the mass transfer rate must be sufficiently large that accretion leads to non-degenerate shell burning on the surface of the white dwarf in order to avoid classical nova explosions that eject the accreted matter and, probably, some of the white dwarf material as well (see, e.g., Nomoto 1982; Iben 1982; Fujimoto 1982; Shen \\& Bildsten 2008; and references therein). This constraint requires the progenitor to be bright and hot, qualities exhibited by the super-soft X-ray sources (SSS; van den Heuvel et al. 1992; Kahabka \\& van den Heuvel 1997). The problem is that there are not enough SSS seen in either spiral or ellliptical galaxies to account for the rate of production of SN~Ia by about a factor of order 100 (Di Stefano 2010a). Similar constraints arise for the DD model. Binary synthesis models require that the progenitor systems go through a phase of rapid accretion onto the primary white dwarf prior to the common envelope phase that reveals the second white dwarf. One way to avoid these constraints is to shroud the progenitor systems in sufficient material that soft X-rays may be produced, but absorbed and transmuted into other wavelengths rather than radiated directly. One possibility is the production of winds from the surface of the accreting white dwarf (Hachisu, Kato \\& Nomoto 1992; Kato \\& Hachisu 1999; Hachisu, Kato \\& Nomoto 2010). A constraint on this particular suggestion is the lack of evidence for such a wind in the remnant of SN Ia 1572 (Badenes et al. 2007).  Here we explore the general constraints on circumstellar matter that might produce sufficient absorption to suppress super-soft ($\\sim$ 0.1 keV) X-rays, describe two lines of evidence that such circumstellar absorption exists, discusss the bands in which such absorbed soft flux might be re-emitted and the liklihood that such systems could be observed. ",
        "conclusions": "Small amounts of CSM matter local to the progenitor binary system of SN~Ia could easily suppress X-ray emission from nuclear burning on the surface of an accreting white dwarf. This suggests that the paucity of SSS to account for the required rate of explosion of SN~Ia in either the SD or DD models is not beyond understanding in terms of local extinction.  Gaining some perspective on the role of SSS in the production of SN~Ia does not address all the issues associated with understanding the progenitors of SN Ia. In the SD model, the mass transfer rate is required to be sufficiently high to avoid degenerate, unstable shell ignition and explosion on the surface of the white dwarf. Published models that satisfy this constraint and also provide a reasonable number of progenitor systems, locally extincted or not, require the mass-transferring secondary star to be a moderately massive main sequence star ($> 1.16$ \\msun; Schaefer \\& Pagnotta 2012) a sub-giant or giant star. The recent advent of SN 2011fe, an apparently normal ``plain vanilla\" SN~Ia, has provided new constraints on the progenitor systems. Nugent et al. (2012) argue that lack of light-curve contamination implies that the secondary star was not a red giant, and more likely to be a main sequence star. Li et al. (2012) use archival images to put limits on the companion and rule out luminous red giants and almost all helium star models. Bloom et al. (2012) show that the exploding star was a white dwarf, as expected, and that the secondary star was likely to have had a radius less than 0.1 that of the Sun, excluding companion red-giant and main-sequence stars that fill their Roche lobes.  SNR 0509-67.5 in the LMC was established by scattered, time-delayed spectra to be a SN~Ia of the SN~1991T spectral subclass that exploded about 400 years ago (Rest et al. 2008). Schaefer \\& Pagnota examined deep HST images of this remnant to put even tighter limits on the progenitor of this explosion. They found that any secondary star must be dimmer than $M_V \\sim 8.4$ mag, ruling out basically all published SD models, including those with companion main sequence stars of greater than about 1 \\msun, sub-giants, giants, and those involving the stripped cores of evolved stars. While one might adopt the dodge that this was a single event responsible for a somewhat peculiar and ill-understood sub-class of SN~Ia, and hence not typical of ``plain vanilla\" SN~Ia, these limits remain a very tight constraint on SD models. Either SD models must be rejected for this system, or some means must be found to impeach the current set of SD models, virtually all of which are based on one-dimensional, spherically-symmetric, non-rotating, non-magnetic accretion that is undoubtedly incorrect, at least in detail.  The bulk of this note was written in January, 2012, independent of the recent posting by Nielsen et al. (2012), who make some of the same points from a different perspective."
    },
    "1208/1208.0543_arXiv.txt": {
        "abstract": "*{We present a cluster analysis of the bright main-sequence and faint pre--main-sequence stellar populations of a field $\\sim 90 \\times 90$\\,pc centered on the \\hii\\ region NGC~346/N66 in the Small Magellanic Cloud, from imaging with HST/ACS. We extend our earlier analysis on the stellar cluster population in the region to characterize the structuring behavior of young stars in the region as a whole with the use of stellar density maps interpreted through techniques designed for the study of the ISM structuring. In particular, we demonstrate with Cartwrigth \\& Whitworth's \\Q\\ parameter, dendrograms, and the $\\Delta$-variance wavelet transform technique that the young stellar populations in the region NGC~346/N66 are highly hierarchically clustered, in agreement with other regions in the Magellanic Clouds observed with HST. The origin of this hierarchy is currently under investigation.}  \\abstract{We present a cluster analysis of the bright main-sequence and faint pre--main-sequence stellar populations of a field $\\sim 90 \\times 90$\\,pc centered on the \\hii\\ region NGC~346/N66 in the Small Magellanic Cloud, from imaging with HST/ACS. We extend our earlier analysis on the stellar cluster population in the region to characterize the structuring behavior of young stars in the region as a whole with the use of stellar density maps interpreted through techniques designed for the study of the ISM structuring. In particular, we demonstrate with Cartwrigth \\& Whitworth's \\Q\\ parameter, dendrograms, and the $\\Delta$-variance wavelet transform technique that the young stellar populations in the region NGC~346/N66 are highly hierarchically clustered, in agreement with other regions in the Magellanic Clouds observed with HST. The origin of this hierarchy is currently under investigation. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.2894_arXiv.txt": {
        "abstract": "A galaxy's mean metallicity is usually closely correlated with its luminosity and mass.  Consequently the most metal-poor galaxies in the local universe are dwarf galaxies.  Blue compact dwarfs and tidal dwarfs tend to deviate from the metallicity-luminosity relation by being too metal-poor or too metal-rich for their luminosity, respectively.  A less pronounced offset separates dwarf spheroidal (dSph) and dwarf irregular galaxies, making the former too metal-rich for their luminosity, which indicates different formation conditions for these two types of dwarfs.  While environment (photo-evaporation through local re-ionization by massive galaxies, tidal and ram pressure stripping) govern the observed morphology-distance relation, intrinsic properties (in particular total mass) play a decisive role in dwarf galaxy evolution with respect to the time and duration of star formation and the amount of enrichment.  The metallicity distribution functions of nearby dwarfs can be understood taking pre-enrichment, gas infall, and winds into account.  Many dwarfs show evidence for inhomogeneous, localized enrichment.  Ultra-faint dSphs, which may have formed their metal-poor stars at high redshift via H$_2$ cooling, show an overabundance of metal-deficient stars as compared to the (inner) Galactic halo, but may, along with classical dSphs, have contributed significantly to the build-up of the outer halo.  The abundance ratios measured in the irregular Large Magellanic Cloud are consistent with the postulated early accretion of irregulars to form the inner Galactic halo. ",
        "introduction": "Most galaxies in the local universe follow a metallicity-luminosity (or metallicity-mass) relation in the sense that more luminous galaxies also tend to be more metal-rich.  Galaxies in the high-metallicity, high-mass regime with current stellar masses greater than $10^{10}$~M$_{\\odot}$ show very little increase in metallicity towards higher masses though \\cite{Panter2008}, while there is a pronounced decline in metallicity towards lower luminosities (or stellar masses) and an increased scatter in the relation \\cite{Panter2008}.  Taking three parameters into account, namely the stellar mass, the gas-phase metallicity, and the star formation rate, galaxies up to a redshift of $z \\sim 2.5$ follow a fairly tight ``fundamental metallicity relation'' in this three-dimensional space \\cite{Mannucci2010}.  At higher redshifts, galaxies show lower metallicities leading to increasing deviations from the relation. Here dilution by infall of pristine gas becomes a dominant factor despite active star formation as long as the dynamical time scales remain shorter than the chemical enrichment time scales \\cite{Mannucci2010}.  At lower redshift outflows of enriched material become more important and a balance between infall and outflows develops \\cite{Mannucci2010}, whereas interacting and merging galaxies deviate from the stellar mass-metallicity trend \\cite{Michel-Dansac2008}.  In this review, we consider metal-poor galaxies in the {\\em local} Universe, i.e., galaxies at the low-luminosity, low-mass end of the metallicity-luminosity relation --- thus, dwarf galaxies.  Exactly where to draw the line between ``dwarf'' and ``giant'' galaxies is more a matter of definition than of physical difference.  Often a simple luminosity criterion is used (e.g., a luminosity $\\leq 0.1\\, L_{\\star}$ or $M_V \\leq -18$ mag; see, e.g., \\cite{Grebel2000}), which also excludes the massive galaxies along the classical Hubble sequence.  Dwarf galaxies cover a variety of morphological types, luminosities, present-day star formation activity, gas content, as well as environments \\cite{Grebel2001}.  Late-type, gas-rich, star-forming dwarfs (such as dwarf spirals, dwarf irregulars, or blue compact dwarfs) are usually found in low-density environments like the outer regions of galaxy clusters, groups, or in isolation in the field. Early-type, gas-deficient, quiescent dwarfs (such as dwarf ellipticals, dwarf spheroidals, or ultra-compact dwarfs) are primarily found in high-density environments, e.g., in the vicinity of massive galaxies in groups and clusters.  This morphology-distance or gas fraction-distance relation \\cite{Grebel2003, Grebel2005} of the dwarf galaxies perpetuates the morphology-density relation observed for giant early-type and late-type galaxies (e.g., \\cite{Oemler1974, Dressler1980}) and suggests that environmental effects may play an important role in shaping the evolution of dwarf galaxies (e.g., \\cite{vandenBergh1994, Grebel2003}).  The best-studied dwarfs are the dwarf galaxies of the Local Group, which are sufficiently close be resolved into individual stars and to even permit us detailed studies of their stellar content down to their oldest main-sequence turn-offs (see \\cite{Monelli2010a, Monelli2010b, Hidalgo2011} for recent examples).  All nearby dwarf galaxies studied in detail so far, regardless of their morphological type, contain very old stars with ages $> 10$ Gyr \\cite{Grebel2004, Glatt2008a, DaCosta2010}.  Though the fractions of these old populations vary and more massive dwarf galaxies usually show more extended episodes of star formation, old populations (i.e., early star formation) appear to be ubiquitous in dwarfs \\cite{Grebel2004}.  In low-mass galaxies the younger and/or metal-rich populations tend to be more centrally concentrated (e.g., \\cite{Grebel1997, Hurley-Keller1999, Zaritsky2000, Harbeck2001, Crnojevic2010, Lianou2010, Lisker2006a, Kniazev2009, Crnojevic2011a, Kirby2011a, Weisz2011}).  Similarly, more metal-rich, younger populations tend to be dynamically colder than older, more metal-poor populations (e.g., \\cite{Tolstoy2004, Battaglia2006, Ibata2006, Battaglia2011}), though not all dwarf galaxies show such population gradients (e.g., \\cite{Koch2007a, Koch2007b, Ural2010}).  Taking advantage of the wealth of information available for nearby, low-metallicity dwarf galaxies and their identifiable, metal-poor old populations, they will be in the focus of this review (although we will also consider galaxies beyond the Local Group).  Many of the Local Group dwarf galaxies are even close enough for detailed spectroscopic abundance analyses of individual member stars, revealing not only the overall metal content, but the star-to-star variations in light and heavy elements of different nucleosynthetic origin (e.g., \\cite{Aden2011, Aoki2009, Cohen2009, Cohen2010, deBoer2012, Hill2000, Kalirai2009, Kirby2011a, Kirby2011b, Koch2006, Koch2007c, Koch2008a, Koch2008b, Norris2008, Sadakane2004, Shetrone2001, Shetrone2003}), which in turn allow us to draw inferences about the conditions under which these metal-deficient galaxies and their stars formed. ",
        "conclusions": ""
    },
    "1208/1208.3341_arXiv.txt": {
        "abstract": "Dynamo action in fully convective stars is a debated issue that also questions our understanding of magnetic field generation in partly convective Sun-like stars. During the past few years, spectropolarimetric observations have demonstrated that fully convective objects are able to trigger strong large-scale and long-lived magnetic fields. We present here the first spectropolarimetric study of a sample of active late M dwarfs (M5-M8) carried out with ESPaDOnS@CFHT. It reveals the co-existence of two distinct types of magnetism among stars having similar masses and rotation rates. A possible explanation for this unexpected discovery is the existence of two dynamo branches in this parameter regime, we discuss here the possible identification with the weak \\textit{vs} strong field bistability predicted for the geodynamo. ",
        "introduction": "In  cool stars, which possess a convective envelope, magnetism is thought to be constantly regenerated against ohmic decay by dynamo effect.  For Sun-like stars the interface layer between the inner radiative zone and the outer convective envelope is generally thought to play a major role in the dynamo processes \\cite[see \\eg][]{Charbonneau10}. Since fully-convective stars -- either main sequence stars below 0.35~\\msun\\ (\\ie with spectral type later than $\\sim$ M4) or young T Tauri stars -- do not possess such an interface layer, generation of magnetic field in their interiors is often thought to rely on a non-solar-type dynamo. However, the precise mechanism and the properties of the resulting magnetic have been a debated issue \\cite[][]{Durney93, Chabrier06, Dobler06, Browning08}.  Two main complementary approaches are successfully applied to study magnetic fields close to the fully-convective transition. On the one hand, by modelling Zeeman broadening of photospheric spectral lines it is possible to assess the magnetic field averaged over the visible stellar disc \\cite[\\eg][]{Reiners06}. This method is therefore able to probe magnetic fields regardless of their complexity but provides very little information about the field geometry. On the other hand, the Zeeman-Doppler imaging technique models the evolution of polarization in spectral lines during at least one rotation period in order to reconstruct a map of the large-scale component of the vector magnetic field on the stellar photosphere.  Spectropolarimetric studies of a sample of M0--M4 dwarfs, conducted with ESPaDOnS and NARVAL, have revealed for the first time a strong change in large-scale magnetic topologies occurring close to the fully-convective boundary. Stars more massive than 0.5~\\msun\\ exhibit large-scale fields of moderate intensity featuring a significant toroidal component and a strongly non-axisymmetric poloidal component, with evolution happening on a timescale of less than 1~yr \\cite[][D08]{Donati08b}. For those in the range 0.25--0.50~\\msun\\ much stronger large-scale fields are observed, which are dominated by the axial dipolar component and show only very limited evolution over successive years \\cite[][M08a,b]{Morin08a, Morin08b}. Comparisons of these large-scale magnetic field measurements with X-ray activity indices or with measurements of the total magnetic field (\\ie at all spatial scales) derived from the analysis of Zeeman broadening of FeH molecular lines, suggest that fully-convective stars are much more efficient at generating large-scale magnetic field than partly-convective ones \\cite[D08,][]{Reiners09b}. ",
        "conclusions": "We present here the main results of the first spectropolarimetric analysis of a sample of active late M dwarfs \\cite[more throughly detailed in][]{Morin10a}. In particular we report the co-existence of two radically different types of magnetism -- strong and steady dipolar field (SD) as opposed to weaker multipolar field evolving in time (WM) -- for stars with very similar masses and rotation periods. One of the foreseen hypothesis to explain these observations is the genuine existence of two types of dynamo in this parameter regime, \\ie bistability. We show that the weak \\textit{vs} strong field dynamo bistability is a promising frame work. The order of magnitude of the observed magnetic field in stars hosting a strong dipolar field, and more conclusively the typical ratio of large-scale magnetic fields measured in the WM and SD groups of stars are compatible with theoretical expectations. We argue that the weak dependency of the magnetic field on stellar rotation predicted for stars in the strong-field regime cannot be ruled out by existing data and should be further investigated.  We do not make any prediction on the extent of the bistable domain in terms of stellar parameters mass and rotation period, this issue shall be investigated by further theoretical work, and by surveys of activity and magnetism in the ultracool dwarf regime.  A dynamo bistability offers the possibility of hysteretic behaviour. Hence the magnetic properties of a given object depend not only on its present stellar parameters but also on their past evolution.  For instance, for young objects episodes of strong accretion can significantly modify their structure and hence the convective energy available to sustain dynamo action \\cite[][]{Baraffe09} initial differences in rotation periods of young stars could also play a role. Because stellar magnetic fields are central in most physical processes that control the evolution of mass and rotation of young stars \\cite[in particular accretion-ejections processes and star-disc coupling, \\eg][]{Bouvier09,Gregory10}, the confirmation of stellar dynamo bistability could have a huge impact on our understanding of formation and evolution of low mass stars."
    },
    "1208/1208.1344_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\noindent Experimental data recently reported by ATLAS \\cite{ATLAS1,ATLAS2} and CMS \\cite{CMS1,CMS2} are consistent with the discovery of  the  Standard Model (SM) Higgs boson, with  mass around $125-126$  GeV.  Such a light Higgs is in good agreement with the indirect indications derived from electroweak precisions constraints~\\cite{EW} under the hypothesis of negligible contributions of physics beyond the SM. Moreover, no clear signal of non-SM physics has emerged yet from collider searches.  Motivated by this experimental situation, in the present paper we try to answer a simple question: can the SM Higgs be responsible for the cosmological perturbations we observe in the universe? As we will see, the answer is: yes.  One of the basic ideas of modern cosmology is that there was an epoch early in the history of the universe when potential, or vacuum, energy associated to a scalar field, the inflaton, dominated other forms of energy density such as matter or radiation. During such a vacuum-dominated era the scale factor grew (nearly) exponentially in time. During this phase, dubbed inflation \\cite{guth81,Starobinsky, lrreview}, a small,  smooth spatial region of size less than the Hubble radius could grow so large as to easily encompass the comoving volume of the entire presently observable universe. If the universe underwent such a period of rapid expansion, one can understand why the observed universe is so homogeneous and isotropic to high accuracy.  Inflation has also become the dominant paradigm for understanding the initial conditions for the Large Scale Structure (LSS) formation and for Cosmic Microwave Background (CMB) anisotropy. In the inflationary picture, primordial density and gravity-wave fluctuations are created from quantum fluctuations ``redshifted'' out of the horizon during an early period of superluminal expansion of the universe, where they are ``frozen'' \\cite{starob79, muk81,hawking82,starobinsky82,guth82,bardeen83}. Perturbations at the surface of last scattering are observable as temperature anisotropy in the CMB. The last and most impressive confirmation of the inflationary paradigm has been recently provided by the data of the Wilkinson Microwave Anistropy Probe (WMAP) mission which has marked the beginning of the precision era of the CMB measurements in space \\cite{wmap7}.  Despite the simplicity of the inflationary paradigm, the mechanism by which  cosmological adiabatic perturbations are generated  is not yet fully established. In the standard picture, the observed density perturbations are due to fluctuations of the inflaton field itself. When inflation ends, the inflaton oscillates about the minimum of its potential and decays, thereby reheating the universe. As a result of the fluctuations each region of the universe goes through the same history but at slightly different times. The final temperature anisotropies are caused by the fact that inflation lasts different amounts of time in different regions of the universe leading to adiabatic perturbations.  Can the  SM Higgs and its potential be responsible for inflation and, at the same time,  the generation of anisotropies?  The answer is: most probably,  not both \\cite{ry}.  The basic problem is that the requirement of having enough e-folds of inflation requires the SM potential to be flat enough, but  this conflicts with the requirement that quantum fluctuation of the Higgs inflaton should also generate the observed power spectrum of anisotropies. Indeed the height  of the SM potential in its flat region is predicted and cannot be arbitrarily adjusted to be as low as needed. This problem can be, in principle, solved by non-minimal coupling of the SM Higgs scalar field $h$  to the Ricci scalar $R$ of the form $\\xi h^2 R$  \\cite{hi,hi1,Barvinsky}. The effect of this interaction is to flatten the Higgs potential (or any other potential) above the  scale $M_{\\rm Pl}/\\sqrt{\\xi}$  providing a platform for slow-roll inflation. A correct normalization of the spectrum of primordial fluctuations fixes the value of the coupling constant $\\xi$ to be larger than about  $10^4$.  This minimal inflationary scenario faces though a couple of issues. First, perturbative unitarity is violated at some scale lower than $M_{\\rm Pl}$ \\cite{hi2} (see however Ref.~\\cite{sibi}), possibly implying the presence of new degrees of freedom which may change the inflationary dynamics. Secondly,  the scenario requires stability of the potential up to the inflationary scale  $M_{\\rm Pl}/\\sqrt{\\xi}$. With  the LHC indication of a Higgs mass in the range  $m_h = (125-126)$  GeV, this simplest version of Higgs inflation is disfavored as the Higgs potential develops an instability at much smaller scales, unless the top mass is below $\\sim 171$ GeV \\cite{hi3,hi4,hi5,Bezrukov:2012sa}.   These considerations lead us to believe that the SM Higgs field may not be responsible for both driving inflation and generating the cosmological perturbations at the same time. Let us therefore be more modest and drop off the requirement that the SM Higgs potential was responsible for inflation.  Our goal is to show that the SM Higgs can nevertheless  play a role in giving rise to the LSS and CMB anisotropies. Indeed, the standard scenario for the generation of the perturbations, where it is the  same scalar field to drive inflation and to source the perturbations,  is not the only option. For instance, an alternative to the standard scenario is represented by the curvaton mechanism \\cite{Linde:1996gt, curvaton1,LW,curvaton3} where the final curvature perturbation $\\zeta$ is produced from an initial isocurvature perturbation associated to the quantum fluctuations of a light scalar field (other than the inflaton), the curvaton, whose energy density is negligible during inflation. The curvaton isocurvature perturbations are transformed into adiabatic ones when the curvaton decays into radiation much after the end of inflation. Alternatives to the curvaton model are those models characterized by the curvature perturbation being generated by an inhomogeneity in the decay rate \\cite{rate} or the mass \\cite{mass} of the particles responsible for the reheating after inflation. Other opportunities for generating the curvature perturbation occur at the end of inflation \\cite{end} and during preheating \\cite{during}   All these alternative models to generate the cosmological perturbations have their strength in the fact that all scalar fields during a period of de Sitter with a mass smaller than the Hubble rate $H$ during inflation are inevitably quantum-mechanically excited with a final superhorizon flat spectrum. Furthermore, they have in common that the comoving curvature perturbation is generated on superhorizon scales when the isocurvature perturbation, which is  associated to the fluctuations of these light scalar fields different from the inflaton,  is converted into curvature perturbation after (or at the end) of inflation.  In the rest of the paper we will therefore assume that there was an inflationary period of accelerated expansion during the primordial evolutionary stage of the universe.  This de Sitter period, induced by some unspecified vacuum energy,  is characterized by a Hubble rate $H$. We will also assume that the perturbations are generated through the Higgs field. This will allow us to play with two independent parameters, the SM Higgs mass $m_h$ and the Hubble rate $H$.  Also, for simplicity, we assume that the inflaton sector does not alter the SM sector, {\\it e.g.} the Higgs potential (see also Ref.~\\cite{cai} for an alternative idea where the Higgs sector of the SM is minimally coupled to asymptotically safe gravity).  Our  considerations will   have direct observational consequences once one realizes that the Hubble rate parametrizes the amount of tensor perturbations during inflation. During the inflationary epoch, tensor perturbations, as for any other massless scalar field, are quantum-mechanically generated. They can give rise to $B$-modes of polarization of the CMB radiation through Thomson scatterings of the CMB photons off free electrons at last scattering \\cite{pol}. The amplitude of the $B$-modes depends on the amplitude of the gravity waves generated during inflation, which in turn depends on the energy scale at which inflation occured. The tensor-to-scalar power ratio is given by $T/S \\simeq (H/3.0\\times 10^{14}\\, {\\rm GeV})^2$. Current CMB anisotropy data impose the upper bound $T/S\\lsim 0.24$ \\cite{wmapping}. The possibility of detecting gravity waves from inflation via $B$-modes is currently being considered by a number of ground, balloon and space based experiments, included the PLANCK experiment. The decomposition of the CMB polarization into $E$- and $B$-modes requires a full sky data coverage and, as such, is limited by the foreground contaminations. The latter introduce a mixing of the $E$ polarization into   $B$ with the corresponding cosmic variance limitation. PLANCK's expected sensitivity is  about $T/S=0.05$ corresponding to a minimum testable value of    $H\\simeq 6.7\\times 10^{13}$ GeV \\cite{plancktensor}. A detection of the tensor mode would therefore imply that the value of the Hubble rate during inflation is larger than about  $10^{13}$ GeV. As we will see, this will have important implications for the the ideas discussed in this paper  The paper is organized as follows. In Section \\ref{sec:perturbation} we describe some  mechanisms by which the SM Higgs could generate the cosmological perturbation. Being ignorant about the exact mechanism, we try to be as generic as possible. In Section \\ref{sec:results} we  present our results and draw our conclusions. ",
        "conclusions": "\\label{sec:results}  We use 2-loop renormalization group (RG) equations for all SM couplings (gauge, Higgs quartic and top-yukawa couplings), and the pole mass matching scheme for the Higgs and top masses, as given in the Appendix of Ref.~\\cite{hi3}. The numerical solution of these equations allows to obtain the  RG-improved effective potential for the SM Higgs \\cite{2loop}, as a function of input parameters, such as the Higgs and top masses. For the top mass we considered $m_t=173.1\\pm 0.7$ GeV (as in Ref.~\\cite{hi5} by combining Tevatron \\cite{Aaltonen:2012ra} and LHC results), while for the QCD gauge coupling $\\alpha_s(M_Z)=0.1184\\pm 0.0007$ \\cite{alphas}. In all simulations we have set $\\alpha_s$ to its central value; the size of the effect of the variation of $\\alpha_s$ within $1\\sigma$ is comparable with the higher-order corrections (such as 3 loops) we are neglecting.       \\begin{figure}[t] \\centering \\includegraphics[scale=1]{3curves.pdf} \\caption{The solution of the field equation (\\ref{eq:fieldeq}) for $h$, normalized to the initial value $h_{\\rm initial}$ (blue line), the second derivative of the potential along the solution $|{\\rm d}^2 V(h)/{\\rm d} h^2|/(3H^2)$ (red line) and the derivative of the field $|{\\rm d}\\ln h/{\\rm d}\\ln a|$ (brown line), as functions of the scale factor $\\ln a$. We have set the parameters as $m_h=125.5$ GeV, $m_t=173.1$ GeV, $H=10^9$ GeV.} \\label{fig:3curves} \\end{figure}   The Higgs field $h$ is initially placed at $h_{\\rm initial}$ where the second derivative of the potential is such that $|{\\rm d}^2 V(h)/{\\rm d} h^2|/(3 H^2)= 10^{-2}$. Then, the field rolls down the potential according to the equation \\be \\ddot h+3 H \\dot h+ V'(h)=0\\,, \\label{eq:fieldeq} \\ee where the dot refers to derivative with respect to $t=H^{-1}\\ln a$. We have conventionally set the scale factor equal to 1 at the initial point $h_{\\rm intial}$.  In Figure \\ref{fig:3curves} we show an example of a situation where the conditions 1. and 2. in the previous section are met and the mechanism works. In fact, the field is slowly rolling down the potential, keeping the second derivative of the potential small (in units of $H^2$), over a wide range of e-folds. So, in this case it is possible to generate nearly scale-independent isocurvature perturbations of the SM Higgs and, through one of the mechanisms mentioned above,  convert them into the observed amount of curvature perturbations $P_{\\zeta}$.  Next, we repeat the analysis by scanning over the Higgs and top masses and looking for what values of $H$ the same situation arises. Of course, an upper limit on the values of $H$ is set by the instability scale $\\Lambda_{\\rm inst}$ at which the Higgs quartic coupling runs negative. For lower values of $H$, down to $\\sim 1$ TeV, we verified that there always exist solutions satisfying conditions 1  and 2 of the previous section, and generating enough curvature perturbations. The analysis of the region $H\\lesssim 1$ TeV is more delicate  as the Higgs field runs very close to its minimum, but we do not pursue this here (we are more interested in the regime where $H$ is relatively large so that tensor modes could be seen by PLANCK).   \\begin{figure}[t] \\centering \\includegraphics[scale=1]{mhvsH.pdf} \\caption{Upper limits on $H$ as a function of $m_h$. Solid black line refers to central value of $m_t$, while dotted blue and dashed red lines correspond to $1 \\sigma$ and $2\\sigma$ variations of $m_t$, respectively. } \\label{fig:mhvsH} \\end{figure}   In Figure \\ref{fig:mhvsH} we show the resulting upper limits on $H$, corresponding to the instability scale $\\Lambda_{\\rm inst}$, for a range of $m_h$ around the observed value. The different lines refer to $m_t$ at the central value or within 1$\\sigma$ and 2$\\sigma$ variations. The mechanism proposed in this paper works for all values of $H$ below the curves. The current exclusion limit on $H$ from CMB data is $H\\lesssim 1.5\\times 10^{14}$ GeV \\cite{wmapping}, which is dark gray band in Fig.~\\ref{fig:mhvsH}). As already mentioned, PLANCK data will possibly be able to exclude values of the Hubble rate down to $H\\simeq 6.7 \\times 10^{13}$ GeV. This limit is shown as a lighter gray band in Fig.~\\ref{fig:mhvsH}. Let us reiterate that the values of $\\beta_h$, see Eq. (\\ref{aaaa}), is not fixed, but is appropriately computed at every point to insure that the cosmological perturbation is correctly normalized, $P_\\zeta^{1/2}=4.8\\times 10^{-5}$. In other words, Figure  2 provides all the possible values of $H$ where the perturbation may be generated by the SM Higgs.   The information in  Figure \\ref{fig:mhvsH} can be read in two ways. For  given  $m_h$ and  $m_t$, there is a maximum $H$ for which the SM Higgs can generate cosmological perturbations, and which may or may not be in the range to be detected in the near future. On the other hand, for a given value of $H$, the hypothesis that the SM Higgs generate perturbations establishes a correlation between $m_h$ and  $m_t$ which can be tested by particle physics measurements.  For example, if the Higgs mass value is confirmed to be $m_h=125.5$ GeV and $m_t$ and  $\\alpha_s$ are at their central values, our mechanism predicts $H\\lesssim 10^{10}$ GeV and therefore tensor perturbations too small to be detected in the near future. On the other hand, if tensor perturbations will be detected with a given $H$, the mechanism we have proposed makes a prediction for a relation between the Higgs and top masses, and we find \\be (m_h)^{\\textrm{B-mode}}\\simeq 128.0 \\GeV+ 1.3\\left(\\frac{m_t-173.1 \\GeV}{0.7 \\GeV}\\right)  \\GeV + 0.9 \\left(\\frac{H}{10^{15} \\GeV}\\right) \\GeV \\pm \\delta_{\\rm th}\\,, \\label{mhfit} \\ee where $\\delta_{\\rm th}\\sim 2$ GeV is a residual theoretical uncertainty from higher order corrections, we have neglected. The formula (\\ref{mhfit}) is valid for $124 \\GeV\\lesssim m_h\\lesssim 127 \\GeV$ and $6.7 \\times 10^{13} \\GeV\\lesssim H \\lesssim 1.5\\times 10^{14} \\GeV$. The information coming from a more accurate experimental determinations of the SM parameters would then allow to either support or rule out our model. Notice also that the result in Eq.~(\\ref{mhfit}) is crude, and just serves as an illustration. A more careful calculation (also including 3-loop $\\beta$-functions and $\\alpha_s$ variation) would be needed in order to extract a detailed and more complete prediction.    Since the purpose of this paper is to show  whether it is possible to generate cosmological perturbations with the SM Higgs,  the level of accuracy we adopted is enough to conclude that the answer is robustly yes. A more accurate determination of the Higgs effective potential, using 3-loop $\\beta$-functions for gauge, Higgs quartic and top-yukawa couplings \\cite{3loop} and 2-loop pole mass matching conditions (as recently presented in Refs.~\\cite{hi5,Bezrukov:2012sa}), would be desirable, but is beyond the scope of the present brief communication.      In conclusion, we have pointed out the possibility that the  SM Higgs might be  responsible for the inhomogeneities we observe in our universe, both in the CMB and in the LSS, if there was an early stage of accelerated expansion driven by some vacuum energy and whose size we agnostically parametrized with the Hubble rate $H$.  Essentially, the Higgs potential may be flat enough to generate nearly scale-independent isocuravture perturbations which are subsequently converted into the observed adiabatic mode. With the recent ATLAS and CMS results for the   Higgs mass, the Higgs can generate the cosmological perturbation in a wide range of the  Hubble rate during inflation,  $H=(10^{10}- 10^{14})$ GeV, depending on the values of the SM parameters. On the other side, if the forthcoming data from the Planck satellite will present some hints of a $B$-mode in the CMB polarization originated from tensor modes, this will identify a well-defined range of the Higgs mass. We have therefore established a very interesting correlation between collider and cosmological measurements, which makes the  mechanism predictive and falsifiable."
    },
    "1208/1208.4098.txt": {
        "abstract": "We present detailed photo+collisional ionization models and kinematic models of the multi-phase absorbing gas, detected within the {\\it HST}/COS, {\\it HST}/STIS, and Keck/HIRES spectra of the background quasar TON 153, at 104~kpc along the projected minor axis of a star-forming spiral galaxy ($z=0.6610$).  Complementary $g'r'i'Ks$ photometry and stellar population models indicate that the host galaxy is dominated by a $\\sim4$ Gyr stellar population with slightly greater than solar metallicity and has an estimated log$M_{\\ast}=11$ and a log$M_{\\rm vir}=13$.  Photoionization models of the low ionization absorption, ({\\MgI}, {\\SiII}, {\\MgII} and {\\CIII}) which trace the bulk of the hydrogen, constrain the multi-component gas to be cold (log$T=3.8-5.2$) and metal poor ($-1.68\\leq [X/H] \\leq -1.64$).  A lagging halo model reproduces the low ionization absorption kinematics, suggesting gas coupled to the disk angular momentum, consistent with cold accretion mode material in simulations.  The {\\CIV} and {\\OVI} absorption is best modeled in a separate collisionally ionized metal-poor ($-2.50\\leq [X/H] \\leq -1.93$) warm phase with log$T=5.3$.  Although their kinematics are consistent with a wind model, given the $2-2.5$~dex difference between the galaxy stellar metallicity and the absorption metallicity indicates the gas cannot arise from galactic winds. We discuss and conclude that although the quasar sight-line passes along the galaxy minor axis at projected distance of 0.3 virial radii, well inside its virial shock radius, the combination of the relative kinematics, temperatures, and relative metallicities indicated that the multi-phase absorbing gas arises from cold accretion around this massive galaxy.  Our results appear to contradict recent interpretations that absorption probing the projected minor axis of a galaxy is sampling winds.   %Our results also appear to contradict the minuscule cold accretion %cross-section predicted by simulations for a massive galaxy at low %redshift, implying that cold accretion may be more prevalent at low %redshift that predicted by simulations. ",
        "introduction": "Over the last decade, simulations have shown that galaxy evolution is highly dependent on gas accretion occurring via two modes: hot and cold accretion. Current cosmological simulations demonstrate that the majority of gas accreted at early epochs onto galaxies occurs via the cold mode, which has temperatures of $T\\sim10^4-10^5$~K and metallicities of $Z\\lesssim 0.01Z_{\\odot}$. Cold-mode gas is preferentially accreted along cosmic filaments/streams and have high densities and low cooling times providing a large supply of gas penetrating through hot halos surrounding galaxies \\citep{keres05,dekel06,ocvirk08,keres09,brooks09, dekel09,ceverino10,stewart11a,stewart11b,freeke11a,freeke11b,faucher-giguere11}. It is expected that cold accretion should comprise no more than ~7\\% of the total {\\HI} mass density at $z\\sim 1$ \\citep{kacprzak11c}.  It is further expected that cold accretion truncates when the host galaxy mass exceeds $\\sim10^{12}$~M$_{\\odot}$, since infalling gas becomes shock heated to the halo viral temperature ($\\sim 10^6$) and is predicted to dramatically reduce the cold accretion cross-section to a tiny fraction \\citep[e.g.,][]{dekel06,keres05,ocvirk08,dekel09,keres09,stewart11a, brooks09,freeke11a,freeke11b} of the observed halo gas cross-section \\citep{kacprzak08,chen10a}.  However, it is expected that these dense filaments can still survive within hot halos and could provide an efficient means of feeding massive galaxies with pristine gas \\citep[e.g.,][]{keres05}.  The study of absorbing foreground gas detected in background quasar spectra allows us to probe these otherwise unobservable comic filaments and outflows. {\\MgII} absorption is ideal for detecting cold mode and hot mode accretion, wind outflows, etc., since it probes gas with a large range of neutral hydrogen column densities, $10^{16} \\lesssim \\hbox{N(\\HI)} \\lesssim 10^{22}$~{\\cmsq} \\citep{archiveI,weakII} with gas temperature of around 30,000$-$40,000~K and average total hydrogen densities of $\\sim0.1$~atoms~cm$^{-3}$ \\citep{cvc03,ding05}. It has also been thoroughly demonstrated that {\\MgII} absorption is produced within gaseous halos surrounding galaxies and is not produced within the intergalactic medium (IGM) \\citep[see][]{cwc-china}.  Over the last decade, strong {\\MgII} absorption has also been observed to directly trace 100$-$1000~{\\kms} galactic-scale outflows \\citep{tremonti07,weiner09,martin09,rubin10,coil11,martin12} that extend out to at least 50~kpc along the galaxy minor axis \\citep{bordoloi11, bouche11,kacprzak12}. Galactic winds have been observed over a large range of redshifts and detected using a range of ions \\citep[see][ and references therein]{steidel10}. Correlations between galaxy colors and star formation rates with {\\MgII} equivalent widths also indirectly suggest that absorption is produced in outflows \\citep{zibetti07,noterdaeme10,nestor11}.  However, {\\MgII} has been observed infalling \\citep{martin12} into highly inclined galaxies with velocities of 100$-$200~{\\kms} \\citep{rubin11}. This is consistent with \\citet{kacprzak11b} who showed that absorption strength is correlated with the orientation of the galaxy major axis, implying that a significant fraction of weaker {\\MgII} absorption systems are likely accreting toward the galaxy via cold flows. The bimodal azimuthal angle distribution of quasar sight-lines around {\\MgII} absorption selected galaxies also suggests that infall occurs along the projected galaxy major axis \\citep{bouche11,kacprzak12}.  These cold-flow streams likely produce a circumgalactic co-rotating gas component that is predominately infalling towards the galaxy and, in absorption, these structures are expected to have $\\sim 100$~{\\kms} velocity offsets relative to the host galaxy {\\it and} in the same direction of galaxy rotation \\citep{stewart11b}. These models are consistent with previous observations of \\citet{steidel02} and \\citet{kacprzak10a} that show {\\MgII} absorption residing fully to one side of the galaxy systemic velocity and usually aligned with expected galaxy rotation direction, with the absorption essentially mimicking the extension of the galaxy rotation curve out into the halo. We expect low ionization states, such as {\\MgI}, {\\MgII}, {\\SiII}, {\\CII} and {\\CIII} to be ideal for tracing cold mode accretion given metallicity, temperatures and densities expected.  A reliable means of determining the origins of the absorbing gas is to obtain both the host-galaxy and absorption-line metallicity. Absorption-line metallicities for a handful of systems have been determined to range between [M/H] $<-1.8$ to $-1$ while existing near sub-L$^{\\star}$ galaxies that have nearly solar metallicities \\citep{zonak04,chen05, tripp05, cooksey08, kacprzak10b, ribaudo11, thom11}. It is postulated that these extremely low metallicity absorption systems are likely accreting onto their host galaxies and possibly trace cold mode accretion, which is still expected for these sub-L$^{\\star}$ galaxies. In the rare case where absorption-line metallicities are larger than the host galaxy is suggestive that the absorption is probing winds \\citep{peroux11}.  Here we target a particular galaxy that has {\\MgII} absorption consistent with disk-like kinematics, possibly tracing cold accretion. The absorption also contains a separate warm phase as indicated by separate strong {\\CIV} absorption that does not coincide with {\\MgII}. We have obtained supplementary {\\it HST}/COS data in order to determine the physical properties of the gas.  In this paper, we perform kinematic and photo+collisional ionization models of multi-phase absorbing gas obtained from {\\it HST}/COS, {\\it HST}/STIS, and Keck/HIRES, which is associated with star-forming spiral galaxy at $z=0.6610$.  In \\S~\\ref{sec:data} we describe our targeted galaxy and our data. We discuss our absorption-line analysis in \\S~\\ref{sec:anal}. In \\S~\\ref{sec:galresults} we describe the host galaxy properties determined from broad-band photometry and stellar population models.  In \\S~\\ref{sec:absresults} we describe the results of our kinematics and photo+collisional ionization models and the physical properties of the absorbing gas.  In \\S~\\ref{sec:dis}, we discuss the possible origins of the absorption and our concluding remarks are in \\S~\\ref{sec:conclusion}. Throughout we adopt an H$_{\\rm 0}=70$~\\kms Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$ cosmology. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we present detailed photo+collisional ionization models and kinematics models of the multi-phase absorbing gas, detected within the {\\it HST}/COS, {\\it HST}/STIS, and Keck/HIRES spectra of the background quasar TON 153, associated with star-forming spiral galaxy at $z=0.6610$. The sightline probes the projected minor axis of the galaxy at projected distance of 0.3 virial radii, well inside the virial shock radius predicted for a galaxy of this mass, implying that if the gas is infalling that it is post shock heated accretion or a cold filament. We obtained followup {\\it HST}/COS data to study other metal-lines in order to determine the halo gas properties and their origins. This galaxy was targeted as a candidate cold accretion probe supported by kinematic and orientation results presented by \\cite{steidel02, kacprzak10a, kacprzak11b}.  Our main results can be summarized as follows:  \\begin{enumerate}   \\item From $g'r'i'Ks$ photometry and stellar population models, we determined that G2 is dominated by a $\\sim4$ Gyr stellar population with slightly greater than solar metallicity abundance and formed at redshift $z\\sim2$. We estimate an $M_{\\ast}=1\\times10^{11}$~M$_{\\odot}$ implying an log $M_{\\rm vir}=12.9$.  \\item The low ionization states, {\\MgI}, {\\SiII}, {\\MgII} and {\\CIII}, have similar absorption kinematics, abundance ratios across the profile, and trace the bulk of the hydrogen, while {\\CIV} and {\\OVI} trace some of the same gas, their kinematics, abundance ratios, and their relative absorption strengths differ.  We infer that the low and high ionization states trace different gas phases.  \\item Modeling the cold gas blueward of G2's systemic velocity, $N(X)^{blue}$, we constrain log$T=3.82-5.23$, $-3.25\\leq$log($U$) $\\leq -3.21$, $-1.87\\leq$ log($n_{_{\\rm H}}$)$ \\leq -1.83$, and $-1.68\\leq [X/H] \\leq -1.64$. The gas is cold and very metal poor: consistent with cold accretion.  We are unable to account for the measured N(\\CIV) and N({\\OVI}) when modeling the cold phase, thus, these ions are likely part of a separate warm collisionally ionized gas phase.  \\item A lagging halo model can account for all of low ionization absorption, hinting that this gas is coupled to the disk and simulations interpret this as a detection cold mode accretion.  \\item Modeling the warm gas blueward of G2's systemic velocity, $N(X)^{blue}$, we find $n_{_{\\rm H}}>-3.5$ and log$T = 5.23-5.29$. Armed with only a conservative limit of hydrogen column density that could be associated with the warm component [$N({\\HI})>$16.8], we estimate $[X/H] \\lesssim -$0.5, although it is highly likely that the metallicity is much lower.  \\item Modeling the warm gas redward of G2's systemic velocity, $N(X)^{red}$, we find hot and metal poor gas with $T=$185,000~K, $-2.50\\leq [X/H] \\leq -1.93$ and $n_{_{\\rm H}} > -3.3$.  \\item The quasar line-of-sight passes along G2's minor axis and a wind model can account for the observed {\\CIV} and {\\OVI} redward and blueward of the galaxy systemic velocity. However, given the $2-2.5$ order of magnitude difference between the galaxy stellar metallicity and the absorption metallicity demonstrates the gas can not arise from galactic winds. \\end{enumerate}  It remains plausible that this low metallicity gas arises from unidentified satellites around the host galaxy or from the incomplete mixing between metal enriched and metal poor halo gas, However, the combination of the relative kinematics, temperatures, and relative metallicities allows us to conclude that the multi-phase gas detected in absorption likely arises from cold accretion around this massive galaxy.  For high mass galaxies the cold accretion cross-section is expected to be a few percent, so our absorption system and others cited in the literature could be a by-chance low probability intersection of a filament, or the resolution effects in the simulations \\citep[see][]{freeke11b} are underestimating the covering fraction of cold flows.  This system also contradicts current results that predict that all absorption detected in quasars probing gas along the projected minor axis of galaxies is produced by winds \\citep{bordoloi11, bouche11,kacprzak12}: This is clearly not the case here.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1208/1208.1491_arXiv.txt": {
        "abstract": "We analyze the density field of 264,283 galaxies observed by the Sloan Digital Sky Survey (SDSS)-III Baryon Oscillation Spectroscopic Survey (BOSS) and included in the SDSS data release nine (DR9). In total, the SDSS DR9 BOSS data includes spectroscopic redshifts for over 400,000 galaxies spread over a footprint of more than 3,000 deg$^2$. We measure the power spectrum of these galaxies with redshifts $0.43 < z < 0.7$ in order to constrain the amount of local non-Gaussianity, $f_{\\mathrm{NL}}^{\\mathrm{local}}$, in the primordial density field, paying particular attention to the impact of systematic uncertainties. The BOSS galaxy density field is systematically affected by the local stellar density and this influences the ability to accurately measure $f_{\\mathrm{NL}}^{\\mathrm{local}}$. In the absence of any correction, we find (erroneously) that the probability that $f_{\\mathrm{NL}}^{\\mathrm{local}}$ is greater than zero, $P(f_{\\mathrm{NL}}^{\\mathrm{local}}>0)$, is 99.5\\%. After quantifying and correcting for the systematic bias and including the added uncertainty, we find -45 $< f_{\\mathrm{NL}}^{\\mathrm{local}} < $ 195 at 95\\% confidence, and $P(f_{\\mathrm{NL}}^{\\mathrm{local}}>0) = 91.0\\%$. A more conservative approach assumes that we have only learned the $k$-dependence of the systematic bias and allows any amplitude for the systematic correction; we find that the systematic effect is not fully degenerate with that of $f_{\\mathrm{NL}}^{\\mathrm{local}}$, and we determine that -82 $< f_{\\mathrm{NL}}^{\\mathrm{local}} < $178 (at 95\\% confidence) and $P(f_{\\mathrm{NL}}^{\\mathrm{local}}>0) = 68\\%$. This analysis demonstrates the importance of accounting for the impact of Galactic foregrounds on $f_{\\mathrm{NL}}^{\\mathrm{local}}$ measurements. We outline the methods that account for these systematic biases and uncertainties. We expect our methods to yield robust constraints on $f_{\\mathrm{NL}}^{\\mathrm{local}}$ for both our own and future large-scale-structure investigations. ",
        "introduction": "Measuring the clustering of galaxies on large scales provides an exciting opportunity to test inflationary models. Slow roll, single field inflation is believed to generate a primordial gravitational potential (hereafter potential) that can be described statistically by a (nearly) Gaussian random field \\citep{Bardeen86}. Alternative inflationary models (e.g., multiple field) predict there to be significant non-Gaussian components to the potential (see, e.g., \\citealt{Wands10} for a review). For a Gaussian random field, the amplitude of fluctuations in 3-point configurations is zero. It is therefore convenient to express the degree of non-Gaussianity, $f_{\\mathrm{NL}}$, as a dimensionless ratio between the amplitude of the bispectrum, $B_{\\Phi}(k_1,k_2,k_3)$, and power spectrum, $P_{\\Phi}(k)$, of the fluctuations of the total potential\\footnote{Specifically, Bardeen's gauge-invariant potential, which is equivalent to the Newtonian potential at sub-horizon scales} $\\Phi$: \\begin{equation} f_{\\mathrm{NL}} \\equiv \\frac{B_{\\Phi}(k_1,k_2,k_3)}{2\\left[P_{\\Phi}(k_1)P_{\\Phi}(k_2)+P_{\\Phi}(k_2)P_{\\Phi}(k_3)+P_{\\Phi}(k_1)P_{\\Phi}(k_3)\\right]}. \\label{eq:fnldef} \\end{equation}  One can denote the portion of the potential that can be described as a Gaussian random field as $\\phi$ and assume that $f_{\\mathrm{NL}} $ is a function of the potential (i.e., that it is `local'). To 2nd order, this approach yields \\begin{equation} \\Phi = \\phi + f_{\\mathrm{NL}}^{\\mathrm{local}}(\\phi^2-\\langle\\phi^2\\rangle). \\end{equation} For this simplest treatment and evaluating Eq. \\ref{eq:fnldef} in the limit where $k$-space triangle configurations satisfy $|\\vec{k}_{12}|\\ll |\\vec{k}_{13}|,|\\vec{k}_{23}|$ (known as the ``squeezed'' limit), it has been shown that a bias in the dark matter halo power spectrum proportional to $f_{\\mathrm{NL}}^{\\mathrm{local}}k^{-2}$ would result (\\citealt{Dalal08,Matt08}). \\cite{Komatsu10} and \\cite{Crem11} have shown that, for standard, single-field inflation, the amplitude of the bispectrum in the ``squeezed'' limit becomes vanishingly small, and thus any detected scale-dependent bias at large-scales would disprove all single-field models, subject to the condition that the field starts in the vacuum. Therefore, measurements of the large-scale clustering of galaxies have the potential to yield significant insight into the physical mechanism that drove inflation.  Primordial non-Gaussianity can also be measured from the bispectrum of cosmic microwave background (CMB) anisotropies (see, e.g., \\citealt{Bartolo04,Komatsu10} and references therein) and, in principle, galaxies (see e.g. \\citealt{Sef07,Scocc12}). WMAP7 found $-10 < f_{\\mathrm{NL}}^{\\mathrm{local}} < 74$ to 95\\% confidence \\citep{KomatsuWMAP7}. To date, the reported constraints on $f_{\\mathrm{NL}}^{\\mathrm{local}}$ from galaxy power spectrum measurements have been competitive with those obtained from the CMB bispectrum. \\cite{Slosar08} used a combination galaxy and quasar clustering measurements to find $-29 < f_{\\mathrm{NL}}^{\\mathrm{local}} < 70$ at 95\\% confidence. \\cite{Xia11} analyzed similar measurements based on updated data to decrease the 95\\% confidence interval to $5 < f_{\\mathrm{NL}}^{\\mathrm{local}} <  84$. \\cite{KomatsuWMAP7} combined their CMB results with the \\cite{Slosar08} result to obtain $-5 < f_{\\mathrm{NL}}^{\\mathrm{local}} < 59$ at 95\\% confidence. Current results thus favour positive $f_{\\mathrm{NL}}^{\\mathrm{local}}$, but do not rule out $f_{\\mathrm{NL}}^{\\mathrm{local}} = 0$. \\cite{Giann12} predict that combining CMB data from the {\\it Planck} mission \\citep{Planck} and galaxy data from a Euclid-like \\citep{euclid} survey will reduce the 1$\\sigma$ uncertainty on $f_{\\mathrm{NL}}^{\\mathrm{local}}$ to 3.  A large value of $f_{\\mathrm{NL}}^{\\mathrm{local}}$ implies larger amplitudes in the 2-point clustering of galaxies at large separations than expected in the standard $\\Lambda$CDM paradigm. Recent studies (e.g., \\citealt{Thomas11,Saw11,Nik12}) have indeed found larger than expected clustering amplitudes at large scales, using photometric Sloan Digital Sky Survey (SDSS; \\citealt{SDSS}) data. However, \\cite{imsys} showed that the excess found in \\cite{Thomas11} was due, at least partially, to stellar contamination, and that systematics correlated with the Galaxy (e.g., stellar density and Galactic extinction) may impart spurious clustering signal at large scales. \\cite{Nik12} find their measurements yield $f_{\\mathrm{NL}}^{\\mathrm{local}} = 90\\pm30$ at 68\\% confidence, but suggest this result may be better interpreted as $f_{\\mathrm{NL}}^{\\mathrm{local}} < 120$ at 84\\% confidence, in light of the systematic concerns. General relativistic (GR) corrections are also expected to cause features in the clustering of galaxies at the largest (super-horizon) scales. These effects are expected to be small compared to that of $f_{\\mathrm{NL}}^{\\mathrm{local}}$, as, e.g., \\cite{Maart12} find that the effects of $f_{\\mathrm{NL}}^{\\mathrm{local}}$ on the power spectrum dominate GR corrections for $f_{\\mathrm{NL}}^{\\mathrm{local}} \\gtrsim 10$.  We analyze the SDSS data release nine (DR9; \\citealt{DR9}) Baryon Oscillation Spectroscopic Survey (BOSS; \\citealt{Daw12bosso}) ``CMASS'' sample of galaxies. This sample has the largest effective volume \\citep{Teg98} of any spectroscopic survey and comprises approximately 1/3 of the final (planned) BOSS CMASS sample. The clustering of this data set has been extensively studied (\\citealt{alph,Nuza12,ReidRSD12,SanchezCos12,Toj12RSD}). In particular, the sample has been simulated with 600 mock galaxy catalogs \\citep{Manera12} and the issues identified by \\cite{imsys} have been addressed via the application of an un-biased weighting scheme \\citep{Ross12}. We thus have the tools to robustly investigate the information content in the large scale clustering, even in light of systematic concerns.  Despite having the largest effective volume of any spectroscopic study, the DR9 CMASS sample is smaller than existing photometric redshift samples of galaxies and quasars (existing quasar samples have roughly 20 times the volume). Given that $f_{\\mathrm{NL}}^{\\mathrm{local}}$ imparts a change in the expected clustering measurements that is most pronounced at large scales and is smoothly varying, we should not expect $f_{\\mathrm{NL}}^{\\mathrm{local}}$ constraints from the DR9 CMASS sample to be competitive even with those one could obtain with the SDSS data release eight \\citep{DR8} photometric redshift sample created in \\cite{imsys} (which has approximately 3 times the angular footprint of our data set and was trained using early BOSS redshifts), let alone a photometric quasar sample. However, we expect our study to have best-quantified the impact of systematics, given the analyses of \\cite{Ross12}.  We primarily use measurements of the spherically-averaged power spectrum, $P(k)$, to constrain $f_{\\mathrm{NL}}^{\\mathrm{local}}$, as the scales most affected by local non-Gaussianity are most isolated in $k$-space. In Appendix \\ref{app:xi}, we test for consistency using the spherically averaged correlation function, $\\xi(s)$. The main purpose of this study is to demonstrate how the ability to constrain $f_{\\mathrm{NL}}^{\\mathrm{local}}$ using the galaxy power spectrum depends on the treatment of systematic uncertainties related to stellar density. We describe the observed and simulated data we use in Section \\ref{sec:data} and our analysis methods in Section \\ref{sec:analysis}. Our results, and their dependence on our treatment of potential systematics, are presented in Section \\ref{sec:results}. In Section \\ref{sec:discussion} we discuss our results in the context of current and future $f_{\\mathrm{NL}}^{\\mathrm{local}}$ measurements, and we summarize our conclusions in Section \\ref{sec:con}.  Unless otherwise noted, we assume a flat cosmology with $\\Omega_{m} = 0.285, \\Omega_{b}=0.0459, h=0.70, n_s=0.96$, and $\\sigma_8=0.8$, as are approximately the best-fit values found by \\cite{SanchezCos12} when fitting the full shape of $\\xi(s)$. ",
        "conclusions": "\\label{sec:con} \\noindent$\\bullet$ We have described a method that quantifies both the systematic bias imparted by Galactic foregrounds on clustering measurements {\\it and} its uncertainty, thus allowing one to obtain un-biased $f_{\\mathrm{NL}}^{\\mathrm{local}}$ constraints with realistic error estimates.  \\noindent$\\bullet$ We find no reliable evidence for non-zero $f_{\\mathrm{NL}}^{\\mathrm{local}}$.  \\noindent$\\bullet$ The data show an extreme preference for our fiducial systematic correction (the application of weights for stellar density when calculating the power spectrum, case [i]) compared to the no systematics correction case (ii): the difference in the minimum $\\chi^2$ is 16.5, when fitting for $f_{\\mathrm{NL}}^{\\mathrm{local}}$ in both cases.  \\noindent$\\bullet$ We have shown that the systematic effect of stars on the DR9 CMASS galaxy density field significantly affects the $f_{\\mathrm{NL}}^{\\mathrm{local}}$ constraints, but that the systematic bias has a different scale dependence than the (convolved) scale dependence of the effect of $f_{\\mathrm{NL}}^{\\mathrm{local}}$. Thus, each effect is detectable in the data, and $f_{\\mathrm{NL}}^{\\mathrm{local}}$ constraints can be obtained even when allowing any amplitude of the systematic correction (case iv), resulting in a 17\\% increase in the width of the 95\\% confidence interval.  \\noindent$\\bullet$ We find that the data exhibit a marginal preference for a stronger systematic correction than provided by our fiducial weights for stellar density, as the minimum $\\chi^2$ decreases by 2.9 when the correction for stellar contamination is 45\\% stronger. We find no physical model that produces a similar change in the minimum $\\chi^2$, but also find no additional systematic effect that can explain the need for such a systematic correction.   We encourage all future studies to incorporate systematic uncertainties in a manner similar to that presented here in order to obtain $f_{\\mathrm{NL}}^{\\mathrm{local}}$ constraints that are robust to systematic uncertainties related to the treatment of the data."
    },
    "1208/1208.0032_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec-19:intro}  According to the standard models of particle physics and cosmology, neutrinos (known and hypothesized) are produced, thermalized, and contribute to the total energy density in the early, hot, dense Universe, regulating the early Universe expansion rate.  Indeed, at the time of big bang nucleosynthesis (BBN), the contributions to the energy density from baryons, dark matter, and dark energy are all subdominant to those from the thermal populations of photons, electrons (\\epm pairs), and neutrinos.  Since the abundances of the elements formed during the first few minutes of the evolution of the Universe depend on the competition between the universal expansion rate and the nuclear and weak interaction rates, the very good agreement between the BBN predictions and observations (see, \\eg,~ref.~\\cite{19-Steigman:2007,19-Steigman:2008,19-Steigman:2010,19-Nollett:2011} for reviews and further references) depends crucially on the early Universe thermalization of neutrinos and places restrictions on the presence of too many (or too few) of them or, of too much ``dark radiation\".  At present, BBN and the cosmic microwave background (CMB) radiation provide the only probes of the cosmic neutrino background.  In addition to their contribution to the total energy density, electron neutrinos and antineutrinos play a special role in regulating the production of \\4he, the second most abundance element in the Universe.  An excess of electron neutrinos over electron antineutrinos (lepton asymmetry;  neutrino degeneracy) or, vice versa, will change the neutron-to-proton ratio during BBN, modifying, mainly, the BBN-predicted primordial helium abundance.  BBN provides a window on the early evolution of the Universe and a probe of particle physics (neutrino physics and more) beyond the standard model (SM).  The primordial abundances of the elements produced in observationally accessible abundances by BBN (primarily D, \\3he, \\4he, \\7li) depend on three fundamental parameters related to cosmology and particle physics: the baryon abundance (related to the Universal baryon asymmetry), the expansion rate of the Universe at BBN (a probe of dark radiation and the cosmic neutrino background) and, any neutrino degeneracy (lepton asymmetry).  \\subsection{Baryon Density Parameter}  The most obvious of these parameters is related to the abundance of the reactants, the baryons (nucleons).  Although the very early Universe may have begun symmetric between matter and antimatter ($n_{\\rm B} = n_{\\rm {\\bar B}}$), long before BBN some yet to be determined mechanism involving the interplay between particle physics (violation of the conservation of baryon number, violation of C and CP symmetries) and cosmology (out of equilibrium evolution) led to a small but crucial local asymmetry between the amount of matter and antimatter in the Universe.  After nucleon-antinucleon annihilation, the excess (nucleons, by definition) survives ($n_{\\rm B} - n_{\\rm \\bar{B}} \\rightarrow n_{\\rm B} \\equiv n_{\\rm N}$) and the number of nucleons in a comoving volume is preserved up to the present epoch (and far into the future as well).  Since the nuclear reaction rates depend on the nucleon density, which decreases as the Universe expands, it is convenient to normalize the nucleon density to the photon density.  After \\epm annihilation, the ratio of the nucleon number density to the photon number density is unchanged as the Universe expands and cools\\footnote{The number of nucleons in a comoving volume is conserved.   Entropy conservation guarantees that, after \\epm annihilation, the number of photons in a comoving volume is also conserved.}.  BBN depends on the baryon density parameter $\\eta_{10}$, defined by \\begin{equation} \\label{eq-19:1} \\eta_{10} \\equiv 10^{10}\\eta_{\\rm B} \\equiv 10^{10}(n_{\\rm B}/n_{\\gamma}). \\end{equation} The present value ($t = t_{0}$, when the photon (CMB) temperature is $T_{0} = 2.725$\\,K) of the baryon density is often measured by comparing the nucleon mass density to the critical mass density ($\\Omega_{\\rm B} \\equiv (\\rho_{\\rm B}/\\rho_{crit})_{0}$) and, the critical mass density depends on the present value of the Hubble parameter, the Hubble constant ($H_{0} \\equiv 100\\,h\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$) \\cite{19-Steigman:2006}, \\begin{equation} \\label{eq-19:2} \\Omega_{\\rm B}h^{2} = \\eta_{10}/273.9. \\end{equation}  Predicting the baryon asymmetry of the Universe is one of the key challenges confronting the search for new physics beyond the standard model.  BBN constraints on $\\eta_{\\rm B}$ can help to identify potentially successful models of new physics.  \\subsection{Expansion Rate Parameter}  The scale factor, $a = a(t)$, describes the evolution of the expansion of the Universe.  During the early evolution of the Universe the expansion rate, as measured by the Hubble parameter, $H \\equiv (1/a)da/dt$, is determined by the total energy density which, during those epochs, is dominated by the contributions from massless or extremely relativistic particles, ``radiation\" (R). \\begin{equation} \\label{eq-19:3} H^{2} = 8\\pi G\\rho/3, \\end{equation} where $G$ is Newton's gravitational constant and $\\rho = \\rho_{\\rm R}$.  New physics may lead to $\\rho_{\\rm R} \\rightarrow \\rho'_{\\rm R}$ (dark radiation) or, to a modification of the cosmology (general relativity) $G \\rightarrow G'$, replacing the SM expansion rate with $H \\rightarrow H' \\equiv SH$.  The expansion rate factor, $S$, quantifies any departure from the standard models of particle physics and/or cosmology.  Prior to the start of BBN and prior to \\epm annihilation (\\eg~$m_{e} \\la T \\ll m_{\\mu}$) the only relativistic SM particles present are the photons (with $g_{\\gamma} = 2$ degrees of freedom or helicities), the \\epm pairs ($g_{e} = 4$), and the N$_{\\nu} = 3$, left-handed neutrinos and their right-handed antineutrinos ($g_{\\nu} = 2$N$_{\\nu}$), so that $\\rho_{\\rm R} = \\rho_{\\gamma} + \\rho_{e} + \\rho_{\\nu}$.  The evolution of the Universe can be scaled out by comparing the total energy density to the energy density in the CMB photons.  Prior to \\epm annihilation, $T_{\\gamma} = T_{e} = T_{\\nu}$, so that accounting for the different contributions to $\\rho_{\\rm R}$ from relativistic fermions and bosons, \\begin{equation} \\label{eq-19:4} {\\rho_{\\rm R} \\over \\rho_{\\gamma}} = 1 + {\\rho_{e} \\over \\rho_{\\gamma}} + {\\rm N}_{\\nu}\\bigg({\\rho_{\\nu} \\over \\rho_{\\gamma}}\\bigg) = 1 + {7 \\over 8}\\bigg[\\bigg({4 \\over 2}\\bigg) + \\bigg({3\\times 2 \\over 2}\\bigg)\\bigg] = {43 \\over 8}, \\end{equation} for N$_{\\nu} = 3$.  The contribution from possible dark radiation (\\eg~sterile neutrinos) may be expressed in terms of an equivalent number of SM neutrinos, \\Deln~\\cite{19-SSG:1977}.  At BBN, which begins prior to \\epm annihilation, N$_{\\nu}$ = 3 + \\Deln.  In this case \\begin{equation} \\label{eq-19:5} \\rho'_{\\rm R} \\equiv \\rho_{\\rm R} + \\Delta{\\rm N}_{\\nu}\\rho_{\\nu}, \\end{equation} or \\begin{equation} \\label{eq-19:6} {\\rho'_{\\rm R} \\over \\rho_{\\gamma}} = {43 \\over 8} + {7 \\over 8}\\Delta{\\rm N}_{\\nu} = {43 \\over 8}\\bigg(1 + {7\\Delta{\\rm N}_{\\nu} \\over 43}\\bigg). \\end{equation} Allowing for dark radiation, the expansion rate factor, $S$, is directly related to \\Deln, \\begin{equation} \\label{eq-19:7} S \\equiv {H' \\over H} = \\bigg({\\rho'_{\\rm R} \\over \\rho_{\\rm R}}\\bigg)^{1/2} = \\bigg(1 + {7\\Delta{\\rm N}_{\\nu} \\over 43}\\bigg)^{1/2}. \\end{equation}  It should be kept in mind that new physics ($S_{\\rm BBN} \\neq 1$) may manifest itself as $G_{\\rm BBN} \\neq G_{0}$ instead of $\\Delta{\\rm N}_{\\nu} \\neq 0$.  In this case, comparing $G_{\\rm BBN}$ when $T \\ga m_{e}$ to its present value, \\begin{equation} \\label{eq-19:8} G_{\\rm BBN}/G_{0} = S_{\\rm BBN}^{2} = 1 + 0.163\\Delta{\\rm N}_{\\nu}. \\end{equation}  After \\epm annihilation the only relativistic SM particles present are the photons and the neutrinos.  The SM neutrinos decouple prior to \\epm annihilation, when $T \\sim 2-3$\\,MeV, so that when the \\epm pairs annihilate, the photons are heated relative to the neutrinos.  On the assumption that the neutrinos are fully decoupled at \\epm annihilation, $T_{\\nu}/T_{\\gamma} = (4/11)^{1/3}$ and, for the SM (\\Deln~= 0), \\begin{equation} \\label{eq-19:9} {\\rho_{\\rm R} \\over \\rho_{\\gamma}} = 1 + \\bigg({\\rho_{\\nu} \\over \\rho_{\\gamma}}\\bigg) = 1 + {21 \\over 8}\\bigg({T_{\\nu} \\over T_{\\gamma}}\\bigg)^{4/3} = 1 + {21 \\over 8}\\bigg({4 \\over 11}\\bigg)^{4/3} = 1.681. \\end{equation} However, in the presence of dark radiation or, ``equivalent neutrinos\" (decoupled, with $T = T_{\\nu} \\neq T_{\\gamma}$), \\begin{equation} \\label{eq-19:10} S^{2} = {\\rho'_{\\rm R} \\over \\rho_{\\rm R}} = 1 + \\bigg({1 \\over 1.681}\\bigg){7 \\over 8}\\bigg({4 \\over 11}\\bigg)^{4/3}\\Delta{\\rm N}_{\\nu} = 1 + 0.135\\Delta{\\rm N}_{\\nu}. \\end{equation} Since the SM neutrinos aren't fully decoupled at \\epm annihilation, they do share some of the energy (entropy) when the \\epm pairs annihilate \\cite{19-Mangano:2005}. This has the effect of increasing the relative contribution of the neutrinos to the total radiation density so that after \\epm annihilation, N$_{\\nu} = 3 + \\Delta{\\rm N}_{\\nu} \\rightarrow$\\,N$_{eff} = 3.046 + \\Delta{\\rm N}_{\\nu}$.  As a result, later in the evolution of the Universe (\\eg~at recombination), $\\rho_{\\rm R}/\\rho_{\\gamma} \\rightarrow 1.692$ and $\\rho'_{\\rm R}/\\rho_{\\gamma} \\rightarrow 1.692 + 0.227\\Delta{\\rm N}_{\\nu}$, so that for $T \\ll m_{e}$, \\begin{equation} \\label{eq-19:11} S^{2} = {\\rho'_{\\rm R} \\over \\rho_{\\rm R}} = 1 + \\bigg({1 \\over 1.692}\\bigg){7 \\over 8}\\bigg({4 \\over 11}\\bigg)^{4/3}\\Delta{\\rm N}_{\\nu} = 1 + 0.134\\Delta{\\rm N}_{\\nu}. \\end{equation} Of course, this post-BBN relation between the expansion rate ($S$) and the equivalent number of neutrinos (\\Deln) is only relevant for those epochs when the Universe is radiation dominated.  BBN codes track the evolution of $S$ from $T \\ga m_{e}$, prior to \\epm annihilation, to $T \\ll m_{e}$, well after \\epm annihilation has ended.  Since it is important for BBN to follow the evolution of the neutron to proton ratio beginning when $T \\ga$~few MeV, prior to \\epm annihilation, \\begin{equation} \\label{eq-19:12} S_{\\rm BBN} \\equiv (1 + 7\\Delta{\\rm N}_{\\nu}/43)^{1/2} = (1 + 0.163\\Delta{\\rm N}_{\\nu})^{1/2}. \\end{equation} A BBN constraint  on $S$ is equivalent to one on \\Deln~(or, on the ratio of $G_{\\rm BBN}$ to its present value, $G_{0}$) and, later in the evolution of the Universe, N$_{eff} = 3.046 + \\Delta{\\rm N}_{\\nu}$.  A BBN determination that $\\Delta{\\rm N}_{\\nu}$ differs from zero at a significant level of confidence can provide evidence for new physics (dark radiation) such as the existence of one, or more, sterile neutrinos (thermally populated) or, a modification of the equations describing the expansion rate of the early Universe ($S_{\\rm BBN} \\neq 1$).  \\subsection{Neutrino Degeneracy Parameter}  Since the charge neutrality of the Universe ensures that any electron excess is tied to the proton excess (the baryon asymmetry), a non-zero lepton asymmetry much larger than the baryon asymmetry ($\\eta_{\\rm B} \\la 10^{-9}$) must be hidden in the neutrino sector.  An excess of neutrinos over antineutrinos (or, vice-versa) requires a non-zero neutrino chemical potential, $\\mu_{\\nu}$.  The dimensionless degeneracy parameter is the ratio of the neutrino chemical potential to the neutrino temperature, $\\xi_{\\nu} \\equiv \\mu_{\\nu}/T_{\\nu}$; $\\xi_{\\nu}$ is preserved as the Universe expands and cools.  In analogy with the parameterization of the baryon asymmetry by $\\eta_{\\rm B} \\equiv (n_{\\rm B} - n_{\\rm {\\bar B}})/n_{\\gamma} \\rightarrow n_{\\rm B}/n_{\\gamma}$, a lepton (neutrino) asymmetry may be parameterized by \\begin{equation} \\label{eq-19:13} \\eta_{\\rm L} = \\eta_{\\nu} = \\Sigma_{\\alpha}\\,{(n_{\\nu} - n_{\\bar{\\nu}})_{\\alpha} \\over n_{\\gamma}} = {\\pi^{3} \\over 12\\zeta(3)}\\Sigma_{\\alpha}\\,\\bigg[\\bigg({\\xi_{\\alpha} \\over \\pi}\\bigg) +  \\bigg({\\xi_{\\alpha} \\over \\pi}\\bigg)^{3}\\bigg], \\end{equation} where the sum is over the three SM neutrino flavors ($\\alpha = e, \\mu, \\tau$).  Generally, mixing among the SM neutrinos ensures that the three chemical potentials are equilibrated.  In the following it is assumed that $\\xi \\equiv \\xi_{e} = \\xi_{\\mu} = \\xi_{\\tau}$.  In this case, \\begin{equation} \\label{eq-19:14} \\eta_{\\rm L} = \\eta_{\\nu} = {\\pi^{3} \\over 4\\zeta(3)}\\bigg({\\xi \\over \\pi}\\bigg)\\bigg[1 +  \\bigg({\\xi \\over \\pi}\\bigg)^{2}\\bigg]. \\end{equation}  An asymmetry between electron neutrinos and electron antineutrinos has a direct effect on BBN through the charged current weak interactions which regulate the neutron-to-proton ratio ($p + e^{-} \\leftrightarrow n + \\nu_{e}, n + e^{+} \\leftrightarrow p + \\bar{\\nu}_{e}, n \\leftrightarrow p + e^{-} + \\bar{\\nu}_{e}$) (see, \\eg~\\cite{19-Beaudet:1976,19-Beaudet:1977,19-Boesgaard:1985,19-Kang:1992,19-Barger:2003,19-Kneller:2004,19-Simha:2008} and further references therein).  Since the relic abundance of \\4he depends directly on the neutron-to-proton ratio when BBN begins (and during BBN), it provides a sensitive probe of any lepton asymmetry.  The abundances of the other light nuclides produced during BBN are less sensitive to $\\xi$.  A subdominant effect (usually) of a non-negligible neutrino degeneracy ($\\eta_{\\rm L} \\gg \\eta_{\\rm B}$) is to enhance to the contribution of the neutrinos to the early Universe energy density.  This is equivalent to a contribution to \\Deln~where, for $\\xi_{e} = \\xi_{\\mu} = \\xi_{\\tau} \\equiv \\xi$, \\begin{equation} \\label{eq-19:15} \\Delta{\\rm N}_{\\nu}(\\xi) = {90 \\over 7}\\bigg({\\xi \\over \\pi}\\bigg)^{2}\\bigg[1 + {1 \\over 2}\\bigg({\\xi \\over \\pi}\\bigg)^{2}\\bigg]. \\end{equation} Note that for $|\\xi| \\la 0.1$, $\\Delta{\\rm N}_{\\nu}(\\xi) \\la 0.013$, which is likely small compared with anticipated uncertainties in \\Deln~inferred from BBN or the CMB.  At present and, likely for the foreseeable future, BBN provides the only window to a universal lepton asymmetry. ",
        "conclusions": "The simplest (least interesting?) assumption is that of no new physics; the standard model with no dark radiation (\\eg~no sterile neutrinos) or a significant lepton asymmetry ($\\eta_{\\rm B} \\ll \\eta_{\\rm L} \\ll 1$).  In this case the SBBN-predicted baryon abundance (see \\S\\,\\ref{sec-19:sm}), $\\Omega_{\\rm B}h^{2} = 0.0218 \\pm 0.0010$, is in excellent agreement with the more precise value found from the CMB, $\\Omega_{\\rm B}h^{2} = 0.0226 \\pm 0.0004$ \\cite{19-Komatsu:2011}.  For this value of the baryon density (and \\Deln~$= \\xi = 0$) the SBBN-predicted abundance of \\3he is consistent with the primordial value inferred from Galactic observations and, the relic abundance of \\4he agrees with the observationally-inferred value adopted here, within $\\sim 1.5\\,\\sigma$.  As is by now well established, the predicted relic lithium abundance exceeds the values of the lithium abundance inferred from observations of the most metal-poor stars in the Galaxy by factor of $\\sim 3 - 4$.  Setting aside for the moment the possibility of a large neutrino degeneracy ($\\eta_{\\nu} \\gg \\eta_{\\rm B}$) but, allowing for the presence of dark radiation (\\Deln~$\\neq 0$), BBN along with the adopted primordial abundances of D and \\4he may be used to constrain \\Deln~and $\\Omega_{\\rm B}h^{2}$.  In this case (see \\S\\,\\ref{sec-19:etannu}) it is found that $\\Delta{\\rm N}_{\\nu} = 0.66^{+0.47}_{-0.45}~({\\rm N}_{eff} = 3.71^{+0.47}_{-0.45}$) and $\\Omega_{\\rm B}h^{2} = 0.0229 \\pm 0.0012$, in excellent agreement with the values of these parameters inferred from various CMB observations \\cite{19-Komatsu:2011,19-ACT:2011,19-SPT:2011,19-SPT:2011a}.  When derived from the CMB, the errors on N$_{eff} = 3.046 + \\Delta{\\rm N}_{\\nu}$ are larger (typically by a factor of $\\sim 1.5 - 2$), and those on $\\Omega_{\\rm B}h^{2}$ smaller (typically by a factor of $\\sim 3$), than the corresponding BBN uncertainties.  This is likely to change when the PLANCK collaboration analyzes its CMB data.  The PLANCK constraint on \\Deln~is expected to be more precise compared to the BBN value by a factor of $\\sim 2.5$, while that on the baryon density parameter should be more precise than the BBN value by an order of magnitude \\cite{19-Hamann:2008,19-Galli:2010}.  The BBN-predicted relic lithium abundance when $\\Delta{\\rm N}_{\\nu} \\neq 0$ is hardly changed from the SBBN case, reinforcing the lithium problem.  As may be seen from Fig.\\,\\ref{fig-19:nnuvseta95}, while the result for \\Deln~(N$_{eff}$) is closer to \\Deln~= 1 than to \\Deln~= 0, offering some support for the existence of one sterile neutrino, it is consistent with \\Deln~= 0 at $\\sim 1.5\\,\\sigma$.  In contrast, the existence of two sterile neutrinos is disfavored by the BBN data at $\\ga 95\\%$ confidence.  \\begin{figure}[h!!] \\begin{center} \\vspace{1cm} \\includegraphics*[scale=0.5]{neffvsomeg12ee.pdf} \\caption{Comparing the BBN predictions of N$_{eff}$ and $\\Omega_{\\rm B}h^{2}$ with those from various CMB determinations: BBN D + \\4he (red filled triangle), BBN D + WMAP7 \\cite{19-Komatsu:2011} $\\Omega_{\\rm B}h^{2}$ (red open triangle), WMAP7 \\cite{19-Komatsu:2011} (blue filled square), ACT \\cite{19-ACT:2011} (green filled pentagon), SPT \\cite{19-SPT:2011} (purple filled circle), SPT + Clusters \\cite{19-SPT:2011a} (purple open circle). \\label{fig-19:neffvsomeg}} \\end{center} \\end{figure}  In Fig.\\,\\ref{fig-19:neffvsomeg} the BBN constraints on \\Deln~and $\\Omega_{\\rm B}h^{2}$ are compared with those from recent CMB analyses.  Here, too, it appears that current data have a preference for one sterile neutrino while being slightly inconsistent with two sterile neutrinos (\\eg~BBN and SPT in Fig.\\,\\ref{fig-19:neffvsomeg}) or with no dark radiation (\\eg~WMAP7 and ACT in Fig.\\,\\ref{fig-19:neffvsomeg}).  Unlike the CMB, BBN has the potential to probe a non-zero (albeit relatively large) lepton asymmetry (neutrino degeneracy).  Current data, driven by the adopted \\4he abundance, are consistent with a small, negative value for the neutrino degeneracy parameter, $\\xi = -0.038 \\pm 0.026$ which, however, is only $\\sim 1.5\\,\\sigma$ from zero.  A more precise result will only come when (if) there is a reduction in the error of the observationally-inferred value of \\Yp, whose uncertainty is dominated by systematics.  \\subsection{Sensitivity Of \\Deln~And $\\xi$ To Primordial Helium And Its Uncertainty} \\label{sec-19:hevstime}  \\begin{figure}[h!!] \\begin{center} \\vspace{1cm} \\includegraphics*[scale=0.4]{hevstime12.pdf} \\hspace{.5cm} \\includegraphics*[scale=0.4]{nnuvstime12a.pdf} \\hspace{.5cm} \\caption{The left panel shows a history of the primordial helium mass fraction (\\Yp) determinations as a function of time.  The same symbols/colors correspond to determinations from collaborations involving many of the same participants and/or the same observational data.  The right panel shows the corresponding chronology of BBN-determined values of \\Deln.  The dashed line shows the SM result, \\Deln~= 0. \\label{fig-19:hevstime}} \\end{center} \\end{figure}  The BBN-predicted helium abundance is sensitive to the early Universe expansion rate ($S$ or, dark radiation \\Deln) and to a lepton asymmetry ($\\xi$) and very insensitive to the baryon density ($\\eta_{\\rm B}$).  The results which have been presented here for \\Deln~and $\\xi$ are mainly driven by the adopted value for the primordial helium abundance, and its uncertainty.  But, the observationally inferred helium abundance, \\Yp, is a quantity which has changed dramatically over time as more and better data have been acquired and more careful analyses of the data have been performed.  In the left panel of Fig.\\,\\ref{fig-19:hevstime} is shown a chronology, over the past $\\sim 20$ years, of the published observational determinations of the primordial helium mass fraction, revealing a nearly monotonic increase of \\Yp~with time.  In the right panel of Fig.\\,\\ref{fig-19:hevstime} the chronology of the corresponding \\Deln~values is shown, mirroring the increase in \\Yp.  Notice that only very recently, within the past 5 -- 7 years, do the data begin to favor $\\Delta{\\rm N}_{\\nu} > 0$.   \\subsection{Constraints On \\Deln~From BBN D And The CMB-Inferred Baryon Density} \\label{sec-19:bbncmb}  The extreme sensitivity of the BBN-inferred estimates of \\Deln~to the adopted helium abundance (and its large errors), is responsible for the relatively large error in the BBN-inferred value of \\Deln.  An alternate approach avoiding \\4he has been suggested by Nollett and Holder (2011) \\cite{19-Nollett:2011} (see, also, Pettini \\& Cooke (2012) \\cite{19-Pettini:2012}).  In the best of all worlds the BBN-inferred parameter values should be compared with those inferred, independently, from the CMB, complemented when necessary to break degeneracies among the parameters by other astrophysical data from, \\eg~large scale structure, supernovae, and the Hubble constant.  In the presence of possible new physics, this would enable a probe of the constancy (or not) of these parameters in the early Universe epochs from BBN until recombination.  However, if it is {\\it assumed} that \\Deln~and $\\eta_{10}$ are unchanged from BBN to recombination, the information provided by BBN using the helium abundance may be replaced with that from the CMB-determined baryon density: $\\eta_{10}({\\rm CMB}) = 6.190 \\pm 0.115$ \\cite{19-Komatsu:2011}.  Using this value in combination with deuterium, $\\eta_{\\rm D} = \\eta_{10} - 6(S - 1) = 5.96 \\pm 0.28$ (for $\\xi = 0$), leads to a smaller estimate of \\Deln~but, with a larger uncertainty resulting from the much weaker dependence of $\\eta_{\\rm D}$ on \\Deln: $\\Delta{\\rm N}_{\\nu} = 0.48^{+0.66}_{-0.63}$ (N$_{eff} = 3.53^{+0.66}_{-0.63}$)\\footnote{For the single, most precise deuterium abundance found by Pettini \\& Cooke \\cite{19-Pettini:2012}, $\\eta_{\\rm D} = 6.10 \\pm 0.24$.  If this abundance is identified with the primordial deuterium abundance, $\\Delta{\\rm N}_{\\nu} = 0.18 \\pm 0.55$ (N$_{eff} = 3.22 \\pm 0.55$).}.  The 68\\% and 95\\% contours for these results are shown in Fig.\\,\\ref{fig-19:cmb95}.  While this provides some support, once again, for the presence of one sterile neutrino, the absence of dark radiation ($\\Delta{\\rm N}_{\\nu} = 0$) is consistent with these results at the $\\sim 68\\%$ confidence level.  For these values of $\\eta_{\\rm D}$ and $\\eta_{10}$ the BBN-predicted helium abundance is Y$_{\\rm P} = 0.2541 \\pm 0.0081$, very close to(well within the errors of) the observationally-inferred value adopted here.  Once again, the BBN-predicted lithium abundance is a problem: A(Li) $= 2.69 \\pm 0.05$.  \\begin{figure}[h!!] \\begin{center} \\vspace{0.75cm} \\includegraphics*[scale=0.5]{nnuvseta95bb.pdf} \\caption{The 68\\% (solid) and 95\\% (dashed) contours in the \\Deln~-- $\\eta_{10}$~plane derived from BBN deuterium and the CMB constraint on the baryon density \\cite{19-Komatsu:2011}. \\label{fig-19:cmb95}} \\end{center} \\end{figure}  This approach, replacing \\4he with the CMB determined baryon density parameter, could also be used to constrain a lepton asymmetry.  The corresponding constraint on the neutrino degeneracy, $\\xi = -0.18 \\pm 0.24$, while entirely consistent with $\\xi = 0$, has an uninterestingly large uncertainty resulting from the very weak dependence of $\\eta_{\\rm D}$ on $\\xi$.  \\subsection{Supplementing BBN With The CMB To Constrain $\\xi \\neq 0$ And $\\Delta{\\rm N}_{\\nu} \\neq 0$} \\label{sec-19:xunnu}  \\begin{figure}[h!!] \\begin{center} \\vspace{1cm} \\includegraphics*[scale=0.5]{xivsnnu12a.pdf} \\caption{The $\\pm 1\\,\\sigma$ band (red) in the $\\xi - \\Delta{\\rm N}_{\\nu}$ plane from BBN using the D and \\4he constraints.  The blue band is the $\\pm 1\\,\\sigma$ range for \\Deln~from Joudaki (2012)~\\cite{19-Joudaki:2012}. \\label{fig-19:xivsnnu}} \\end{center} \\end{figure}  The BBN-predicted primordial light element abundances depend on all three of the key parameters \\{$\\eta_{10}, \\Delta{\\rm N}_{\\nu}, \\xi$\\}.  However, the uncertainty in the observationally-inferred relic abundance of \\3he, along with the lithium problem(s), leaves only two, relatively well constrained primordial abundances, those for D and \\4he.  From BBN alone and these abundances, all three parameters can't be determined independently but, one of them can be eliminated resulting in a relation (degeneracy) between the remaining two.  For example, $\\eta_{\\rm D}$ and $\\eta_{\\rm He}$ may be used to eliminate $\\eta_{10}$, leading to $\\xi = \\xi(\\Delta{\\rm N}_{\\nu};\\,y_{\\rm DP},{\\rm Y_{\\rm P}})$, where \\begin{equation} 145\\xi = 106(S - 1) + \\eta_{\\rm D} - \\eta_{\\rm He}. \\label{eq-19:xivsnnudhe} \\end{equation} This constraint on $\\xi$ versus \\Deln~is shown by the red band in Fig.\\,\\ref{fig-19:xivsnnu}.  Without an independent constraint on \\Deln, the degeneracy between $\\xi$ and \\Deln~seen in Fig.\\,\\ref{fig-19:xivsnnu} cannot be broken.  The CMB provides such a constraint.  If BBN, using the observationally-inferred D and \\4he abundances, is now supplemented with an independent constraint on \\Deln~from the CMB, then a combined constraint, allowing for both dark radiation and lepton asymmetry, may be found.  Using the CMB results from WMAP7 \\cite{19-Komatsu:2011}, ACT \\cite{19-ACT:2011}, and the SPT \\cite{19-SPT:2011,19-SPT:2011a}, along with complementary constraints from large scale structure, the Hubble constant, supernovae, and galaxy clusters, Joudaki (2012) \\cite{19-Joudaki:2012} finds N$_{eff} = 3.87 \\pm 0.42$, corresponding to $S = 1.065 \\pm 0.032$.  Using this result in Eq.\\,\\ref{eq-19:xivsnnudhe} leads to $\\xi = 0.009 \\pm 0.035$, entirely consistent with $\\xi = 0$.  The $2\\,\\sigma$ upper bound to $|\\xi|$ is 0.079, corresponding to an upper bound to $\\Delta{\\rm N}_{\\nu}(\\xi) = 0.008$.  Once again, the BBN-predicted lithium abundance is high, A(Li)~$= 2.70 \\pm 0.06$, reinforcing the lithium problem.  A similar approach may be used to eliminate \\Deln~instead of $\\eta_{10}$ to find $\\xi = \\xi(\\eta_{10};\\,y_{\\rm DP},{\\rm Y_{\\rm P}})$ and, to use the CMB for a constraint on $\\eta_{10}$.  But this approach, which is also consistent with $\\xi = 0$, leads to a less precise constraint: $\\xi = -0.012 \\pm 0.052$.     The BBN results presented here provide modest support for the presence of dark radiation ($\\Delta{\\rm N}_{\\nu} \\ga 0$ at $\\sim 1.5\\,\\sigma$), to go along with the more robust evidence for dark matter and dark energy.  The current observational data (D and \\4he), while allowing for the presence of a sterile neutrino (\\Deln~= 1 at $\\la 1\\,\\sigma$), disfavors the presence of two sterile neutrinos ($\\Delta{\\rm N}_{\\nu} = 2$ at $\\sim 2.8\\,\\sigma$).  If, instead, it is {\\it assumed} that \\Deln~= 0, a small, but non-zero lepton asymmetry is favored, also at the $\\sim 1.5\\,\\sigma$ confidence level.  In contrast, if both \\Deln~and $\\xi$ are allowed to vary freely and, if BBN (D and \\4he) is supplemented by a CMB constraint on \\Deln, a vanishing lepton asymmetry ($\\xi = 0.009 \\pm 0.035$) is favored.  Currently, as may be seen from Fig.\\,\\ref{fig-19:neffvsomeg}, there is very good agreement between the BBN and CMB constraints on the baryon density and dark radiation (the CMB is insensitive to a small or modest lepton asymmetry).  For many years BBN provided the best constraints on the baryon density ($\\eta_{10}~{\\rm or}~\\Omega_{\\rm B}h^{2}$) and on dark radiation (\\Deln~or $S$), as well as the only constraint on lepton asymmetry.  With WMAP7 \\cite{19-Komatsu:2011} and other CMB datasets \\cite{19-ACT:2011,19-SPT:2011,19-SPT:2011a} the best constraints on the baryon density now are from the CMB, which allow for a factor of $\\sim 2 - 3$ more precise determination of $\\eta_{10}$.  However, in the present, pre-PLANCK era, BBN still provides the best dark radiation constraint, albeit with an uncertainty smaller than that from the CMB by only a factor of $\\sim 1.5 - 2$.  It is expected that with the publication of the PLANCK data the dark radiation torch will pass to the CMB.  Depending on what PLANCK finds, it may be possible to establish the presence of dark radiation (a sterile neutrino?) at the $\\sim 5\\,\\sigma$ level {\\it if}, for example, PLANCK should find, \\Deln~$= 1 \\pm 0.2$.  If, however, PLANCK should find (the best or worst of all worlds?) \\Deln~$= 0.5 \\pm 0.2$, the presence dark radiation will be favored but, that of a sterile neutrino will be somewhat disfavored\\footnote{It is perhaps worth recalling that the contribution of a light, thermalized scalar corresponds to \\Deln~= 4/7 = 0.57.}.  For all the possibilities considered here (\\Deln~$= \\xi = 0$; $\\xi = 0$, $\\Delta{\\rm N}_{\\nu} \\neq 0$; $\\Delta{\\rm N}_{\\nu} = 0$, $\\xi \\neq 0$; $\\Delta{\\rm N}_{\\nu} \\neq 0$, $\\xi \\neq 0$), the BBN-predicted lithium abundance hardly changed at all ($2.68 \\pm 0.06 \\leq {\\rm A(Li)} \\leq 2.70 \\pm 0.06$).  This insensitivity is easy to understand since the BBN-predicted lithium and deuterium abundances are strongly correlated \\begin{equation} \\eta_{\\rm Li} = \\eta_{\\rm D} + 3[(S - 1) - \\xi]. \\label{eq-19:livsd} \\end{equation} Because the values of $3(S - 1)$ and $3\\xi$ are almost always small compared to $\\eta_{\\rm D}$, the corrections to a perfect lithium -- deuterium correlation are generally at only the few percent level.  A solution to the lithium problem is not to be found with dark radiation or a lepton asymmetry.  \\subsection{Anticipating The Future}  The future for the key parameters related to the baryon abundance ($\\eta_{\\rm B}$) and the presence, or not, of dark radiation (\\Deln) lies with the CMB and the anticipated results from the PLANCK mission. While it is impossible to predict the central values PLANCK will find for $\\eta_{\\rm B}$ or \\Deln, it is possible to forecast the precision to be expected from the PLANCK data and analyses \\cite{19-Hamann:2008,19-Galli:2010}.  Such forecasts suggest that the uncertainty in the baryon abundance determination will be of order $\\sigma(\\eta_{10}) \\approx 0.03$, nearly an order of magnitude better than the current BBN precision.  The same forecasts suggest that \\Deln~will be constrained to $\\sigma($\\Deln$) \\approx 0.2$, or better.  This would result in an improvement over the current BBN precision by a factor of $\\sim 2.5$.  If it is {\\it assumed} that the PLANCK values of the key parameters are identical to those at BBN, ignoring their possible evolution from the epoch of BBN until recombination, then these values may be used in combination with BBN to predict the relic abundances\\footnote{In contrast with the current approach of using BBN in concert with the observationally inferred relic abundances to predict the values of the key parameters.}.  For example, for deuterium, it is anticipated that PLANCK will constrain $\\eta_{\\rm D}$ with a precision of $\\sigma(\\eta_{\\rm D}) \\approx 0.1$ or, to $\\la 2\\%$ for $\\eta_{\\rm D} \\approx 6$.  The largest uncertainty in the BBN-predicted deuterium abundance at present and in this anticipated future arises from uncertain nuclear reaction rates \\cite{19-Nollett:2011}.  It can be hoped that this uncertainty may be reduced by new laboratory data, reducing the error in the BBN-predicted value of $y_{\\rm DP}$ by perhaps a factor of $\\sim 2$.  This would lead to a reduction in the error in the inferred value of $\\eta_{\\rm D}$ by nearly a factor of two, $\\sigma(\\eta_{\\rm D}) \\approx 0.3 \\rightarrow  0.15$.  For helium, PLANCK may constrain $\\eta_{\\rm He}$ to $\\sigma(\\eta_{\\rm He}) \\approx 1.6$ which, while still large, is a factor $\\sim 2.3$ smaller than the current BBN uncertainties.  This suggests that using the CMB determined values of $\\eta_{10}$ and \\Deln, the BBN-predicted primordial helium mass fraction will be known to $\\sigma({\\rm Y_{P}}) \\la 0.003$, a precision anticipated to also be attainable in an independent determination of \\Yp~from the CMB \\cite{19-Hamann:2008,19-Galli:2010}.  For lithium, PLANCK may constrain $\\eta_{\\rm Li}$ to $\\sigma(\\eta_{\\rm Li}) \\approx 0.06$ or, to better than $\\sim 1\\%$ for $\\eta_{\\rm Li} \\approx 6$.  However, as for deuterium, the precision of the BBN-predicted primordial lithium abundance is limited by the nuclear physics uncertainties ($\\sim 10\\%$).  Nonetheless, it will be very interesting to see if the PLANCK data support or, possibly eliminate, the lithium problem(s).  Although the CMB is insensitive to a lepton asymmetry, as may be seen from Eq.\\,\\ref{eq-19:xivsnnudhe}, a combination of BBN and CMB constraints on \\Deln, $y_{\\rm DP}$, and \\Yp~can constrain a neutrino degeneracy, provided that the lepton asymmetry is very large compared to the baryon asymmetry.  For example, for the anticipated CMB constraints on $\\sigma(\\Delta{\\rm N}_{\\nu}) \\approx 0.2$ and on the primordial helium abundance, $\\sigma({\\rm Y_{P}}) \\approx 0.003$, along with a BBN constraint on $y_{\\rm DP}$, $\\sigma_{\\xi} \\approx 0.018$ or, $\\sigma_{\\eta_{\\rm L}} \\approx (\\pi^{2}/4\\zeta(3))\\sigma_{\\xi} \\approx 0.036 \\approx 6\\times 10^{7}\\eta_{\\rm B}$.  \\subsection{Summary}  This review finds itself on the cusp of potentially great changes.  Current BBN and CMB data provide strong support for the presence of (at least) three SM neutrinos, thermally populated during the early evolution of the Universe. This provides indirect support for the so far invisible, relic neutrino background.  The new CMB and large scale structure data have the potential to constrain the baryon asymmetry and the presence, or not, of dark radiation to new levels of precision, testing BBN and the current estimates of the relic abundances of the light elements.  It will be of great interest to compare and contrast the current BBN results with those from the new data and to see what we may learn about new physics, including neutrino physics, beyond the standard models of particle physics and cosmology.  \\begin{center} {\\bf Acknowledgments} \\end{center} I am pleased to acknowledge informative conversations and email exchanges with R. Cyburt, M. Fumagalli, K. Nollett, M. Pettini, J.~X. Prochaska.  My research is supported at OSU by the DOE."
    },
    "1208/1208.2037_arXiv.txt": {
        "abstract": "We derive the fundamental parameters (temperature and luminosity) of 107\\,619 \\emph{Hipparcos} stars and place these stars on a true Hertzsprung--Russell diagram. This is achieved by comparing {\\sc BT-Settl} model atmospheres to spectral energy distributions (SEDs) created from \\emph{Hipparcos}, \\emph{Tycho}, SDSS, DENIS, 2MASS, \\emph{MSX}, \\emph{AKARI}, \\emph{IRAS} and \\emph{WISE} data. We also identify and quantify from these SEDs any infrared excesses attributable to circumstellar matter. We compare our results to known types of objects, focussing on the giant branch stars. Giant star dust production (as traced by infrared excess) is found to start in earnest around 680 L$_\\odot$. ",
        "introduction": "\\label{IntroSect}  Spectral energy distributions (SEDs) have long been the primary method of understanding stars. Colour--magnitude diagrams, which can be quickly made from photometric data, enable one to explore various facets of stellar populations, such as stellar mass and evolutioanry state. However, these does not present the information at its most basic physical level: the stellar temperature and luminosity. These represent the fundamental ideals of stellar modelling, and are theoretically free from biases introduced by photometric calibration, interstellar reddening and similar phenomena.  While transformations to these parameters can be achieved through colour--temperature relations and bolometric corrections, these are limited in scope. Most importantly, the wavelength coverage of the observations means that well-defined solutions do not always exist for these relations (e.g. for very red stars, or for observations only covering wavelengths longer than the SED peak). Using the entire wavelength coverage available allows better determination of stellar temperature when a wide temperature range is present among a stellar sample. This also allows finer control of data quality. All-sky surveys are, in particular, prone to contain some poor-quality data due to the large flux range they are required to cover, which leads to the saturation of bright sources, and the volume of data, which limits the ability to match photometric routines to particular situations (e.g.\\ in areas of high stellar density or nebular emission). Stellar variability can also cause improper colours to be reported, which can be reduced by using multiple epochs or, equivalently, multi-wavelength data. In this manner, we can provide more-robust estimates of parameters for individual objects, allowing them to be placed on the true, physical Hertzsprung--Russell (H--R) diagram.  Perhaps the greatest benefit, however, is the ability to detect excess flux at a particular wavelength, by providing a reference model flux against which fluxes in individual photometric filters can be compared. This is particularly helpful in the infrared, where colour--magnitude diagrams based on only part of the SED can fail to identify sources exhibiting emission in addition to the stellar photosphere. Predominantly, these sources are either very young stars (pre-main-sequence T Tauri stars or Herbig Ae/Be stars), rapid rotators (classical Be stars), or evolved stars. This latter group is mostly comprised of mass-losing red and asymptotic giant branch (RGB/AGB) stars, on which we focus our discussion.  Previously, only colour--magnitude diagrams have been used to interpret our closest stellar neighbours (e.g.\\ \\citealt{PLK+95}). We are now able to take the data returned by the \\emph{Hipparcos} satellite \\citep{Perryman97,vanLeeuwen07} and match it with other all-sky surveys to produce a true H--R diagram of the local Solar neighbourhood.  In doing so, we can identify and characterise stars with weak infrared excesses which may be otherwise missed by conventional colour cuts. While this has been attempted previously (\\citealt{IMI+10,Groenewegen12}; we later discuss these papers in context), this work represents the first time such a process has been applied to the entire \\emph{Hipparcos} dataset and in the context of the stars' absolute, fundamental parameters. ",
        "conclusions": "\\label{ConcSect}  In this work, we have demonstrated the use of spectral energy distribution fitting to determine the fundamental parameters of the \\emph{Hipparcos} star sample. We have further used this information to quantify excess flux over the entire optical, and near- and mid-infrared region of each SED. We have combined these excesses to determine those stars showing an excess of infrared flux, and cross-correlated literature identifications to examine the cause of that excess over different regions of the H--R diagram, comparing our results to the key studies of \\citet{IMI+10} and \\citet{Groenewegen12}. We find we cannot reproduce the infrared excess and dust production claimed by the latter paper.  Our analysis has focussed on the \\emph{Hipparcos} data catalogue: data which is now over 20 years old and, despite showing its age, provides the best estimate of distances to nearby stars we have. The launch of \\emph{Gaia}, and the completion of further all-sky surveys such as Pan-STARRS, SDSS and \\emph{WISE}, will allow a similar analysis to be performed on many times more objects. Automated techniques, building on the kind demonstrated here, will be necessary to analyse and classify the objects which come from these surveys, in order to gain a full and comprehensive understanding of our corner of the Galaxy and its inhabitants."
    },
    "1208/1208.0174_arXiv.txt": {
        "abstract": "Gravitational microlensing, when finite size source effects are relevant, provides an unique tool for the study of source star stellar atmospheres through an enhancement of a characteristic polarization signal. This is due to the differential magnification induced during the crossing of the source star. In this paper we consider a specific set of reported highly magnified, both single and binary exoplanetary systems, microlensing events towards the Galactic bulge and evaluate the expected polarization signal. To this purpose, we consider several polarization models which apply to different types of source stars: { hot, late type main sequence and cool giants}.  As a result we compute the  polarization signal $P$, which goes up to P=0.04 { percent} for { late type} stars and up to a few percent for cool giants, depending on the underlying physical polarization processes and atmosphere model parameters. Given a $I$ band magnitude at maximum magnification of about  12, and a typical duration of the polarization signal up to 1 day, we conclude that the currently available  technology, in particular the polarimeter in FORS2 on the VLT, potentially may allow the detection of such signals. This observational programme may take advantage of the currently available surveys plus follow up strategy already routinely used for microlensing monitoring towards the Galactic bulge (aimed at the detection of exoplanets). In particular, this allows one to predict in advance for which events and at which exact time the observing resources may be focused to make intensive polarization measurements. ",
        "introduction": "Gravitational microlensing, initially developed to search for MACHOs in { the} Galactic halo and near the Galactic disc \\citep{Pacz86,Macho93,Eros93,Ogle93,Ogle94} has become nowadays a powerful tool to investigate several aspects of stellar astrophysics and also to search for extrasolar planets orbiting around lens stars.  Indeed, microlensing gives the opportunity to study the star's limb-darkening profile, which is the variation of the intensity from the disc center to the limb, and thus to test stellar atmosphere models. At the same time, microlensing leads to the discovery and the detailed characterization of exoplanetary systems when planetary deviations in the light-curves expected for single-lens events are detected (see \\cite{Dominik10} and \\cite{Gaudi10} for recent reviews).  Microlenses can spatially resolve a source star thanks to caustic structures created by a lens system \\citep{SEF}. { Caustics are formed by a set of closed curves,}  along which the point source magnification is formally infinite, with a steep increase in magnification in their vicinity. This increase is so steep that the characteristic length scale of the differential magnification effect is of the order of a fraction of the source star radius. In this way different parts of the source star are magnified by substantially different amounts. The resulting lensing light-curve deviates from the standard form expected for a point source event \\citep{Witt94,Gould94,Alcock97}  and the analysis of the deviations enables to measure the limb-darkening profile of the lensed star.   Early works \\citep{Witt95,Loeb-Sasselov95,Valls-Gabaud98,Heyrovsky03} have pointed out the sensitivity of microlensing light-curves to limb-darkening, with the aim to help to remove the microlensing model parameter degeneracy. The specific use of microlensing as a tool to study stellar atmospheres was proposed later \\citep{Hendry98,Gaudi-Gould99}, in particular to probe atmospheres of red giants in the Galactic bulge \\citep{Heyrovsky00}.  {The best candidate events for studying stellar atmospheres are highly magnified microlensing events, which also show relevant finite size source effects. For these events the lens and the background source star are almost aligned and the lens passes over the surface of the source star ({\\it transit} events). Although relatively rare, these events potentially contain unique information on the stellar atmosphere properties of the source star as shown by \\cite{Fouque10} and \\cite{Zub10}. Indeed, besides the brightness profile of a remote source star disc, highly magnified events with large finite size source effects allow to measure the lens Einstein radius $R_{\\rm E}$ (if the physical radius $R_*$ of the source is known) and provide a unique chance to study spectroscopically Galactic bulge stars.  The light-curve analysis of highly magnified events is also sensitive to the presence of lens planetary companions, in particular when the planet-to star distance is of the order of $R_{\\rm E}$. The same opportunity for studying stellar atmospheres is offered by binary microlensing due to caustic crossing as the source passes through fold \\citep{SEF} and cusp caustics \\citep{Schneider92,Zakharov95}. }  The aim of the present paper is to consider polarization variability of the source star light for real events, taking into account different polarization  mechanisms according to the source star type. Indeed, variations in the polarization curves are similar to finite source effects in microlensing when { color effects } may appear due to limb darkening and color distribution across the disc \\citep{Witt94,Bogdanov95a,Bogdanov95b, Gaudi-Gould99,Bogdanov00,Bogdanov02,Dominik05,Heyrovsky07}.  It is known that the light received from the stellar limb may be significantly (up to about 12 { percent}) polarized due to the so called Sobolev -- Chandrasekhar effect \\citep{Sobolev49,Chandra60,Sobolev63}. Polarization parallel to the limb of the source is caused by photon (Thomson) scattering off free electrons, when the light passes through the stellar atmosphere. This polarization mechanism is effective { for hot stars of any luminosity class} which, indeed, have a free electron atmosphere.  It is also known that, by a minor extent, the continuum spectrum of the {  main sequence stars of late type} is linearly polarized by coherent (Rayleigh) scattering on neutral hydrogen in its ground state and, with a minor contribution, by scattering on free electrons \\citep{Stenflo99}. The polarization in { late type stars} has been measured only for the Sun (for which due to the distance the projected disc is spatially resolved) and also in this case { as for hot stars}, the polarization gets its maximum value near the solar limb, due to the most favorable geometry there.  However, the light received from the stars is usually unpolarized, { since the flux from each stellar disc element is the same.} A net polarization of the stellar light may be introduced by some suitable asymmetry in the stellar disc (e.g. eclipses, tidal distortions, stellar spots, fast rotation, magnetic fields) and also in the propagation through the interstellar medium \\citep{BKS85,Dolginov95}.  In the microlensing context, polarization in the stellar light is induced due to the proper motion of the lens star through the source star disc, during which different parts are magnified by different amounts. Therefore, the gravitational lens scans the disc of the background star giving rise not only to a time dependent gravitational magnification of the source star light but also to a time dependent polarization.  As for { the} limb-darkening brightness profile, the polarization degree is maximized when the source trajectory crosses regions with high magnification gradient. This occurs in {\\it transit} events since the magnification of the source star flux increases as the lens approaches closer to the source star, and during binary microlensing events since now there is the opportunity for caustic crossing.  Accordingly, we consider the polarization variability for the {\\it transit} events (with lenses passing over source stars) selected by \\cite{Choi} and for a subset of exoplanetary events towards the Galactic bulge \\citep{Gaudi10}. { Exoplanetary events are binary lens systems characterized by values of the planet-to-star mass ratio $q \\ll 1$ and have smaller star-to-planet distance $d$ as compared to the separation of the stars in a binary system. A full treatment of the polarization in binary microlensing events will be the subject of a subsequent paper. Here, selecting a representative sample of exoplanetary events, we show how the polarization works in binary microlensing.}  The idea that polarization induced by an electron scattering atmosphere could be enhanced by gravitational microlensing and eventually observed was first raised and investigated in relation to supernovae by \\cite{SW87}. In particular, the observation of a variable polarization  due to a supernova beam expanding in a self-gravitating system constituted by individual masses ($ \\ut > 10^{-3}~M_{\\odot}$) may indicate the presence of dark matter objects.  Then, \\cite{Simmons95a}, \\cite{Simmons95b} and \\cite{BCS96} pointed out that the polarization of star light induced by an electron scattering atmosphere is enhanced by the microlensing effect and that the relative motion of source and lens causes a time variation of the source polarization degree. \\cite{Simmons95a} and \\cite{Simmons95b} also present a numerical calculation of the polarization degree induced by a single-lens (the Schwarzschild lens), showing that during a microlensing event the polarization profile has a double peak for {\\it transit} events, (where a part of the source disc is aligned with the lens and observer), while it has a single peak in the {\\it bypass} events (where the source trajectory remains outside the lens).  Assuming that the source star has an electron scattering atmosphere, \\cite{Agol96} calculated the time-dependent polarization of a star being gravitationally lensed by a binary system. Polarization as high as $P \\simeq 1$ {percent} can be achieved if the star crosses a caustic or passes near a cusp; otherwise, the maximum polarization is $\\simeq 0.1$ { percent}.  Polarization by non-compact microlenses \\citep{Gurevich95,Zakharov96a,Zakharov96b,Zakharov99,Zakharov10} was also analyzed \\citep{Belokurov98,Zakharov98} since a source may cross caustics arising in the model.  The most likely candidates for observing polarization variability during microlensing events would be young, hot giant star sources, because they have electron scattering atmospheres needed for producing limb polarization through Thomson scattering \\citep{Simmons95b}. Unfortunately, the bulge of the Galaxy does not contain a large number of hot giant stars. However, polarization may be also induced by the scattering of star light off atoms, molecules and dust grains in the adsorptive atmospheres of evolved, cool stars as shown by \\cite{Simmons02} and \\cite{Ignace06}. These more ubiquitous stars, that do not have levels of polarization as high as those predicted by the Chandrasekar model, may display an intrinsic polarization of up to several { percent}, due to the presence of stellar winds that give rise to extended adsorptive envelopes. This is the case for many cool giant stars, in particular for the red giants. Such evolved stars constitute a significant fraction of the lensed sources towards the Galactic bulge, the LMC \\citep{Alcock97} and the M31 galaxy \\citep{Sebastiano10}, making them valuable candidates for observing variable polarization during microlensing events.  Polarization measurements on ongoing microlensing events can be useful for further characterizing them, for testing stellar atmosphere models  and to complement finite source measurements \\citep{Gould94}. In this respect, the detection of a variable polarization leads to an independent measure of the angular Einstein radius $R_{\\rm E}$ of the lens, the position angle of the lens and the velocity direction in the sky \\citep{Yoshida}.  Of course, since accurate polarization measurements cannot be obtained with a survey telescope, alert systems are necessary allowing other larger telescopes to take polarimetric measurements during a microlensing event.  For definiteness, we consider the observed sample of highly magnified, single-lens {\\it transit} events \\citep{Choi} and a subset of exoplanetary events observed towards the Galactic bulge \\citep{Gaudi10}. For these events we calculate the polarization profiles as a function of time taking into account the nature of the source stars. As an illustration, we also consider the expected polarization signal for the PA-99-N2 event towards M31 \\citep{Paulin03}. ",
        "conclusions": "The combined effects of large magnification and finite size source effects in some microlensing events may allow to get relatively large values of the polarization of the light from the source stars. In this paper we calculate the polarization profile as a function of the time for a selected sample of both single and exoplanetary microlensing events observed towards the Galactic bulge, by taking into account the nature of the source star: { hot, late type main sequence and cool giant stars.} Indeed, different polarization mechanisms take place in the stellar atmospheres, depending on the source star type: photon (Thomson) scattering on free electrons, coherent (Rayleigh) scattering off atoms and molecules, and photon scattering on dust grains, { for hot, late type and cool giant stars} (with extended atmospheres), respectively.  The analysis of the polarization curves for single-lens, highly magnified microlensing events towards the Galactic bulge has shown that the polarization degree of the stellar light can reach values as high as 0.04 { percent} at the peak in the case of {late type source stars} and up to a few {percent} in the case of cool giant source stars (red giants) with extended envelopes. For these events the primary lens crosses the  source star disc ({\\it transit} events) and relatively large values of $P$ are thereby produced due to large finite source effects and the large magnification gradient throughout the source star disc. The time duration of the peak of the polarization signal may vary, from 1h to 1day, depending on the source star radius and the lens impact parameter.  Similar values of polarization may also be obtained in exoplanetary events when the source star crosses the primary or the planetary caustics. While in the former case (as for single-lens events) the peak of the polarization signal always occurs at symmetrical points with respect to the instant $t_0$ of maximum magnification, in the latter case the polarization signal may occur at any (and generally unpredictable) time during the  event.  The natural question which arises is whether such polarization signals may be detectable with the present or near future technology. Polarimeters are nowadays available on large telescopes and the best possibility for measuring polarization in the R and I bands is offered at present by the polarimeter in  FORS2 on ESO's VLT telescope. With this instrument it is possible to measure the polarization for a 12 mag source star with a precision of 0.1 { percent} in 10 min integration time, and for a 14 mag star in a 1h. Indeed, for a few of the events considered in Tables \\ref{table2} and \\ref{table4} the peak $I_0$ magnitude is $\\simeq 12$ with the expected maximum polarization degree lasting up to 1 day, thus suggesting that polarization measurements in highly magnified microlensing events constitute a realistic target of opportunity for currently available instruments.   It goes without saying that polarization measurements in microlensing events require an alert system able to predict in advance the instant of the occurrence of the polarization peak. A alert system is already in operation based on OGLE and MOA survey data \\footnote{www.OGLE.astrouw.edu.pl and www.phys.canterbury.ac.nz/MOA}, and it is particularly efficient in the case of highly magnified events, allowing to focus observing resources to make intensive observations around the peak of the event.  We  emphasize that polarization measurements in highly magnified microlensing events offer an unique opportunity to probe stellar atmospheres of Galactic bulge stars. Besides the interest related to stellar astrophysics, the analysis  of the polarization profile, which reflects that of the magnification light-curve, given sufficient observational precision, may in principle provide independent constraints on the lensing parameters also for exoplanetary events."
    },
    "1208/1208.2347_arXiv.txt": {
        "abstract": "We perform a detailed study on the dynamics of a relativistic blast wave with the presence of a long-lived reverse shock (RS). Although a short-lived RS has been widely considered, the RS is believed to be long-lived as a consequence of a stratification expected on the ejecta Lorentz factors. The existence of a long-lived RS makes the forward shock (FS) dynamics to deviate from a self-similar Blandford-McKee solution. Employing the ``mechanical model'' that correctly incorporates the energy conservation, we present an accurate solution for both the FS and RS dynamics. We conduct a sophisticated calculation of the afterglow emission. Adopting a Lagrangian description of the blast wave, we keep track of an adiabatic evolution of numerous shells between the FS and RS. An evolution of the electron spectrum is also followed individually for every shell. We then find the FS and RS light curves by integrating over the entire FS and RS shocked regions, respectively. Exploring a total of 20 different ejecta stratifications, we explain in detail how a stratified ejecta affects its blast wave dynamics and afterglow light curves. We show that, while the FS light curves are not sensitive to the ejecta stratifications, the RS light curves exhibit much richer features, including steep declines, plateaus, bumps, re-brightenings, and a variety of temporal decay indices. These distinctive RS features may be observable if the RS has higher values of the micophysics parameters than the FS. We discuss possible applications of our results in understanding the GRB afterglow data. ",
        "introduction": "\\label{section:introduction}   The central engine of a gamma-ray burst (GRB) ejects a relativistic outflow (called an ejecta) with high Lorentz factors. As the ejecta interacts with a surrounding ambient medium, a relativistic blast wave develops. The blast wave consists of two shock waves: the forward shock (FS) wave sweeping up the ambient medium and the reverse shock (RS) wave propagating through the ejecta. The shocked ambient medium is separated from the shocked ejecta by a contact discontinuity (CD), and a compressed hot gas between the FS and RS is called a ``blast''.  Without an extended activity of the central engine, the RS is expected to be short-lived if the ejecta is assumed to have a constant Lorentz factor $\\Gej$. The RS vanishes as it crosses the end of the ejecta. The blast wave then enters a self-similar stage where the FS dynamics is described by the solution of Blandford \\& McKee (1976) (hereafter BM76). This FS emission has been believed to be the main source of the long-lasting, broad-band afterglows (M\\'esz\\'aros \\& Rees 1997; Sari et al. 1998). The short-lived RS emission would be then important only briefly in the early afterglow phase. Thus, it was proposed to explain a brief optical flash detected in some GRBs (M\\'esz\\'aros \\& Rees 1997,1999; Sari \\& Piran 1999a,b). The dynamical evolution of such a short-lived RS with a constant $\\Gej$ was studied analytically (Sari \\& Piran 1995; Kobayashi 2000), under the assumption of an equality of pressure across the blast wave.  However, a general view on the structure of the ejecta should include the possibility that the ejecta emerges with a range of the Lorentz factors. The shells with lower Lorentz factors gradually ``catch up'' with the blast wave as it decelerates to a comparable Lorentz factor. Therefore, the RS wave is believed to be long-lived in general.  An example with a long-lived RS, where a power-law ejecta interacts with a power-law ambient medium, was studied analytically by assuming a constant ratio of two pressures at the FS and RS (Rees \\& M\\'esz\\'aros 1998; Sari \\& M\\'esz\\'aros 2000).  The structure or stratification of the ejecta and ambient medium could be in fact even more general. There is no reason why it should take only a constant or power-law profile. Uhm (2011) (hereafter U11) presented a semi-analytic formulation for this class of general problems where the ejecta and ambient medium can have an arbitrary radial stratification.  U11 takes into account a radial spread-out and spherical expansion of such a stratified ejecta and finds which shell of this evolved ejecta gets passed by the RS at a certain time and radius.  U11 then finds the dynamics of the blast wave with a long-lived RS, by employing two different methods: (1) an equality of pressure across the blast wave (mentioned above) and (2) the ``mechanical model'' (Beloborodov \\& Uhm 2006).  U11 shows that the two methods yield significantly different dynamical evolutions and demonstrates that the method (1) does not satisfy the energy conservation law for an adiabatic blast wave while the method (2) does.  The mechanical model does not assume either an equality of pressure across the blast wave or a constant ratio of two pressures at the FS and RS.  It shows that the ratio of two pressures should in fact evolve in time as the blast wave propagates.  Besides these theoretical considerations, recent early afterglow observations led by {\\it Swift} revealed a perplexing picture regarding the origin of GRB afterglows. In contrast of a simple power-law decay feature as expected from the standard afterglow theory, the X-ray data show more complicated features including initial rapid declines, plateaus, and flares (e.g. Tagaliaferri et al. 2005; Burrows et al. 2005; Nousek et al. 2006; O'Brien et al. 2006; Chincarini et al. 2007) that reveal rich physics in the early afterglow phase (Zhang et al. 2006; Zhang 2007). More puzzlingly, some GRBs show clear chromatic behaviors of the X-ray and optical afterglows (e.g. Panaitescu et al. 2006; Liang et al. 2007, 2008). It is now evident that the FS alone cannot interpret the broad-band afterglow data for the majority of GRBs.   Uhm \\& Beloborodov (2007) and Genet et al. (2007) independently showed that the RS-dominated afterglow flux could reproduce some observed afterglow features, given the assumption that the FS emission is suppressed. In this paper, we study in great detail the dynamics and afterglow light curves of GRB blast waves with a long-lived RS. The purpose is to investigate how different ejecta stratifications affect their blast wave dynamics and afterglow light curves. We explore various types of the ejecta stratification and unveil that there exists a whole new class of the blast wave dynamics with a rapid and strong evolution of the RS strength. In order to find an accurate solution for both the FS and RS dynamics, we make use of U11 with the mechanical model. As explained above, this allows for the blast wave with a long-lived RS to satisfy the energy conservation, by introducing a pressure-gradient across the blast wave region.  We perform a sophisticated calculation of afterglow emission, invoking a Lagrangian description of the blast wave. In the widely-used analytical afterglow model (e.g. Sari et al. 1998), it is assumed that the entire shocked material forms a single zone with same energy density and magnetic field. The electron energy distribution is solved only in the energy space, with no consideration of spatial distribution within the shocked region. Beloborodov (2005) (hereafter B05) described a more sophisticated Lagrangian method, in which the postshock region is resolved into subshells using a Lagrangian mass coordinate. B05 studies an evolution of the magnetic field and power-law spectrum of electrons for each subshell as the blast wave propagates. However, the postshock material in B05 is not resolved in radius and all the subshells are located at the same radius. Also, the pressure and energy density in B05 are assumed to be constant throughout the postshock material. Improving on B05, here we have a spatial resolution into the blast region, allowing for our Lagrangian shells to have their own radius. We further introduce a pressure profile that smoothly varies over the blast. This is because the pressure at the FS differs from the pressure at the RS, as discussed above. As the blast wave propagates, we keep track of an evolution of the pressure, energy density, and adiabatic index of every shell on the blast. We also keep track of an evolution of the magnetic field and power-law spectra of electrons of all shells.  Finally, in order to calculate synchrotron radiation from a spherical shell on the blast, we analytically find an observed spectral flux for a distant observer, taking into account the effects of the shell's radial velocity and spherical curvature. We integrate this flux over the entire blast to find the sum of emissions from all the shells between the FS and RS.  In Section~\\ref{section:dynamics}, we briefly summarize how we find the dynamics of a blast wave with a long-lived RS. In Section~\\ref{section:light_curves}, we describe in detail our method of calculating afterglow light curves. Numerical examples are presented in Section~\\ref{section:numerical_examples}, which exhibit various features on the blast wave dynamics and afterglow light curves. Our results are summarized in Section~\\ref{section:discussion} (Discussion) and Section~\\ref{section:conclusion} (Conclusion). ",
        "conclusions": "\\label{section:conclusion}   We have investigated in detail the dynamics and GRB afterglow light curves for a relativistic blast wave with a long-lived RS. This long-lived RS is formed in a stratified outflow ejected from a central source, which spreads out radially and forms various density structures in the flow due to its stratification. Due to spreading of Lorentz factor, this ejecta flow gradually catches up with the blast wave and adds its kinetic energy into the blast, naturally maintaining a long-lived RS wave. As a result, this blast wave is not in the self-similar stage, as is described in the BM76 solution. Instead, the blast wave is being continuously pushed by the RS, which displays various forms of energy injection scenarios.   In order to find such dynamics of a blast wave with a long-lived RS, we make use of U11 with the mechanical model and perform detailed numerical calculations. Investigating a total of 20 different shapes of ejecta stratifications, we explain the effects and consequences of radial spreadings on the FS and RS dynamics. In particular, we show that there exists a whole new class of the RS dynamics with fast and strong evolutions. The FS dynamics is also shown to exhibit consistent behaviors for those diverse types of energy injections (Figure~\\ref{fig:dyn_all}). A high accuracy shown in Figure~\\ref{fig:dyn_all} indicates that we are presenting here a ``precision dynamics'' for the blast waves with a long-lived RS.   Employing a Lagrangian description of the blast wave, we perform a sophisticated calculation of afterglows. In particular, our calculation has (1) a spatial resolution into the blast wave region and (2) a pressure profile that smoothly varies over the blast. For every shell on the blast, we keep track of an evolution of (1) the thermodynamic quantities of shocked gas (pressure, energy density, adiabatic index, etc) and (2) the magnetic field and power-law spectrum of electrons. The FS and RS light curves are found by integrating over the entire FS and RS shocked regions, respectively, while making use of an analytic expression for observed spectral flux $\\delta F_{\\nuobs}^{\\rm obs}$, which we derive here in terms of an observed frequency $\\nuobs$ and observer time $\\tobs$ (Section~\\ref{section:curvature_effect}).   The resulting afterglow light curves display interesting features. Since the FS strength mainly depends on the Lorentz factor of the blast, the FS light curves do not sensitively depend on the ejecta stratification. The strength of the RS, on the other hand, sensitively depends on the ejecta density flow $\\rhoejRS$ entering the RS, so that rich afterglow features show up in the RS light curves. As demonstrated with our 20 examples, the RS light curves naturally produce diverse and distinctive features (Figure~\\ref{fig:cRX_all}) as we move through the space of ejecta stratifications (Figure~\\ref{fig:gej_all}). In particular, designing proper stratifications in the ejecta, the RS light curves reproduce many observed X-ray features, including various temporal breaks (Figures~\\ref{fig:cRX_1a1b1c} and \\ref{fig:cRX_3a3b3c}) and decay indices (Figures~\\ref{fig:cRX_1a5a5b} and \\ref{fig:cRX_6a6b4d}), plateaus (Figures~\\ref{fig:cRX_2a2b2c}, \\ref{fig:cRX_3a3b3c}, and \\ref{fig:cRX_4a4b4c4d}), steep declines (Figures~\\ref{fig:cRX_2a2b2c}, \\ref{fig:cRX_4a4b4c4d}, and \\ref{fig:cRX_6c6d6e}), bumps (Figure~\\ref{fig:cRX_3a3b3c}), and re-brightenings (Figure~\\ref{fig:cRX_6c6d6e}). Since the FS and RS could have different efficiency in particle acceleration and the GRB ejecta is likely more magnetized than the ambient medium, it is plausible that the RS emission would become as bright as or even brighter than the FS emission. Therefore, we believe that the RS could be a strong candidate to account for the observed GRB afterglows."
    },
    "1208/1208.0342_arXiv.txt": {
        "abstract": "We use a sample of 45 core collapse supernovae detected with the Advanced Camera for Surveys on board the {\\it Hubble Space Telescope} to derive the core collapse supernova rate in the redshift range $0.1<z<1.3$. In redshift bins centered on $\\langle z\\rangle=0.39$,  $\\langle z\\rangle=0.73$, and   $\\langle z\\rangle=1.11$, we find rates 3.00$^{+1.28}_{-0.94}$$^{+1.04}_{-0.57}$, 7.39$^{+1.86}_{-1.52}$$^{+3.20}_{-1.60}$, and 9.57$^{+3.76}_{-2.80}$$^{+4.96}_{-2.80}$, respectively, given in units of  yr$^{-1}$Mpc$^{-3}~10^{-4}~h_{70}^3$.  The rates have been corrected for host galaxy extinction, including supernovae missed in highly dust-enshrouded environments in infrared bright galaxies. The first errors are statistical while the second ones are the estimated systematic errors. We perform a detailed discussion of possible sources of systematic errors and note that these start to dominate over statistical errors at $z>0.5$, emphasizing the need to better control the systematic effects. For example, a better understanding of the amount of dust extinction in the host galaxies and knowledge of the supernova luminosity function, in particular the fraction of faint $M \\gsim-15$~supernovae, is needed to better constrain the rates. When comparing our results with the core collapse supernova rate based on the star formation rate, we find a good agreement, consistent with the supernova rate following the star formation rate, as expected. ",
        "introduction": "Supernovae (SNe) mark the end of the life cycle of certain stars. Studies of SNe are important both for understanding the physics leading to and driving these explosions, as well as the impact that the SNe have on their environments. As the main producers of heavy elements, SNe are pivotal for understanding the chemical evolution in galaxies as well as the intergalactic medium via SN-driven outflows from galaxies. SNe are also thought to be one of the producers of dust in the universe (e.g., review by Kozasa et al. 2009). Furthermore, SNe tie into feedback processes regulating galaxy formation. To understand how these processes evolve with cosmic time, it is important to understand how the number of exploding SNe changes with redshift, i.e., the evolution of the cosmic SN rate (SNR). Since SNe, particularly the core collapse (CC) SNe, have massive progenitors with short main-sequence lifetimes, the rate should closely follow the star formation rate (SFR), offering a direct and independent way of measuring the SFR and at the same time the metal enrichment rate. While this method is straightforward in principle, it is hampered particularly by the need for corrections due to dust extinction in order to derive accurate rates. For Type Ia SNe, determining the rate offers a way to estimate the delay time between the formation of the progenitor star and its explosion as an SN, which may shed light on the still mostly unknown scenarios leading up to the explosion of the degenerate white dwarfs that are assumed to be the progenitors of Type Ia SNe (e.g., Madau et al. 1998; Dahlen \\& Fransson 1999; Gal-Yam \\& Maoz 2004; Maoz \\& Gal-Yam 2004).  Due to their potential as standard candles used to measure the expansion rate of the universe (Riess et al. 1998; Perlmutter et al. 1999; Astier et al. 2006), a large number of surveys aiming at detecting Type Ia SNe have been conducted during the last decade, resulting in a multitude of rate estimates. For a recent compilation of rates from the literature, see Graur et al. (2011).  While these results show a fair consensus that the Type Ia SNR increased by at least a factor of $\\sim$3 between redshift $z=0$~and $z\\sim1$, there is still an uncertainty in the rates at $z\\gsim 1$, reflecting the uncertainty in such searches, the limited statistics and relatively high systematic errors concerning searches at $z>1$. Moreover, optical searches are also hampered by the redshifting of the SN spectral energy distribution (SED), making them drop out of optical filters at $z\\gsim$1. Deep surveys in the near-infrared (IR), such as the {\\it Hubble Space Telescope (HST)} based CANDELS (Grogin et al. 2011; Koekemoer et al. 2011) and CLASH (Postman et al. 2012) programs, should help alleviate this scarcity of high-redshift Type Ia SNe. Rodney et al. (2012) demonstrate the possibility of detecting and follow-up of Type Ia SNe at $z>1.5$~using IR searches.  While Type Ia SNe have been the focus of a large number of surveys during the last decade, there has been a lack of dedicated surveys aiming at finding CC SNe at cosmological distances. Most of the existing surveys suffer from severe selection biases since the foremost aim has been to select clean samples of Type Ia SNe and little follow-up has been devoted to non-Type Ia SN transients, including CC SNe. Furthermore, CC SNe are harder to detect and follow on the light curve because they are typically $\\sim$2 mag fainter than Type Ia SNe. Therefore, there are significantly fewer CC SNRs reported in the literature compared to Type Ia SNRs.  The local CC SNR (outer distances between 11 and $\\sim$200 Mpc) has been determined by Cappellaro et al. (1999), Smartt et al. (2009), Li et al. (2011), Botticella et al. (2012), and Mattila et al. (2012), while at cosmological distances rates have been determined at $z=0.21$~by Botticella et al. (2008), $z=0.26$~by Cappellaro et al. (2005; these rates are, however, superseded by the rates in Botticella et al. 2008), $z=0.3$~by Bazin et al. (2009), $z=0.66$~by Graur et al. (2011), $z=0.39$~and $z=0.73$~by Dahlen et al. (2004), and most recently at $z=0.39$~and $z=0.73$~by Melinder et al. (2012).  Most rates are consistent with each other and they increase with look-back time, reflecting the increased SFR between $z=0$~and $z=1$. However, there are some significant differences in the measured local rates. Furthermore, there are also claims by Horiuchi et al. (2011) that the SNR in general is a factor of $\\sim$2~lower than predicted from the SFRs, suggesting that there are possible unaccounted systematic effects that may affect the determined rates. Most importantly, necessary corrections to account for the dimming of the SNe by dust in the host galaxies are highly uncertain. The effect of dust extinction is also expected to increase with redshift since the rest frame probed in optical searches approaches the UV part of the spectrum where the effect of extinction is more severe. There should also be a population of CC SNe almost completely hidden from optical searches in highly extinguished luminous infrared galaxies (LIRGs) and ultra-luminous infrared galaxies (ULIRGs; see, e.g.,  Mannucci et al. 2007; Mattila et al. 2004, 2012). Since the fraction of the star formation that occurs in U/LIRGs increases significantly with look-back time and actually dominates the total star formation at $z\\sim 1-2$~ (e.g., Magnelli et al. 2011), we expect the number of CC SNe hidden in these environments to increase with redshift. The CC SNe hidden in such galaxies must be accounted for when the aim is to derive the complete rates of CC SNe.  Also, having a good knowledge of how the SN population is divided into different subtypes and the luminosity function of each subtype is essential for deriving accurate rates. In particular, knowing the fraction of faint CC SNe ($M\\gsim$-15) is important when deriving rates since these will remain mostly undetected and must be corrected for when estimating rates.  In this paper, we use the SN sample detected in the extended GOODS/PANS SN survey conducted with $HST$/ACS 2002--2005 to derive the rate of CC SNe to redshift $z\\sim$1.1. In a previous paper (Dahlen et al. 2004), we presented CC SNR to $z$=0.7 based on a subsample from the same search ($\\sim$36\\% of the search that was available at that time). The present paper accompanies Dahlen et al. (2008), where we presented the Type Ia SNR to $z\\sim$1.6 based on the full sample. This paper is organized as follows. In Section 2, we describe the SN search and present the sample. In Section 3, we describe how SNRs are calculated. The results are presented in Section 4 together with a discussion in Section 5. We summarize our results in section 6. Throughout this paper we assume a cosmology with $\\Omega_{\\rm M}$=0.3, $\\Omega_{\\Lambda}$=0.7, and $h$=0.7. Magnitudes are given in the Vega system. ",
        "conclusions": "We have used a sample of 45 CC SNe, out of a total of 62 detected with ACS during the $HST$-extended SN search in Cycles 11--13, to derive the CC SNR to $z\\sim$1.1. Our main conclusions are as follows. \\begin{itemize} \\item{After correcting for host galaxy extinction and the missing fraction in highly obscured environments, we find rates 3.00$^{+1.28}_{-0.94}$$^{+1.04}_{-0.57}$~at $\\langle z\\rangle$=0.39, 7.39$^{+1.86}_{-1.52}$$^{+3.20}_{-1.60}$~at $\\langle z\\rangle$=0.73, and 9.57$^{+3.76}_{-2.80}$$^{+4.96}_{-2.80}$~at $\\langle z\\rangle$=1.11. The rates are given in units of yr$^{-1}$~Mpc$^{-3}~10^{-4}~h_{70}^3$. The first errors represent statistical while the second ones are the estimated systematic errors.} \\item{Our rates at $z\\gsim$0.4 are consistent with those expected from the cosmic SFR.} \\item{Statistical uncertainties dominate our rate estimates in the low-redshift bin, however; systematic errors start to dominate at $z>0.5$.} \\item{While the most recent local CC SNRs are consistent with those expected from the SFR, there is a discrepancy between some earlier derived $z\\lsim$0.3 rates and the SFR.  As noted in Horiuchi et al. (2011) and Mattila et al. (2012), this is likely due to a combination of SN missed due to dust extinction and/or faint intrinsic luminosity and unaccounted systematic effects.} \\end{itemize}"
    },
    "1208/1208.4946_arXiv.txt": {
        "abstract": "The Galaxy and the stars in it form a hierarchical system, such that the properties of individual stars are influenced by those of the Galaxy. Here, an approach is described which uses hierarchical Bayesian models to simultaneously and empirically determine the mean distance-extinction relationship for a sightline and the properties of stars which populate it. By exploiting the hierarchical nature of the problem, the method described is able to achieve significantly improved precision and accuracy with respect to previous 3D extinction mapping techniques.  This method is not tied to any individual survey and could be applied to any observations, or combination of observations available. Furthermore, it is extendible and, in addition, could be employed to study Galactic structure as well as factors such as the initial mass function and star formation history in the Galaxy. ",
        "introduction": "The Milky Way Galaxy occupies a unique role in astronomy. As we are located within it, we are able to observe and analyse it and its constituents in a manner not possible with any other galaxy. However, this also means that we lack a global view of it. Thus, in order to analyse the Galaxy's structure and history we are forced to infer distances to stars. A task considerably complicated by the requirement to disentangle the effects of interstellar extinction.  Recent years have been characterised by a growth in survey astronomy. Wide area photometric surveys such as 2MASS \\citep[][]{Skrutskie_Cutri.2006} and SDSS \\citep[][]{York_Adelman.2000} have revolutionised astronomy on all scales, from brown dwarfs to cosmology. Meanwhile, there has also been significant development in surveys specifically studying the Galactic Plane, which is home to the majority of stars in the Galaxy. Such Galactic plane surveys include IPHAS \\citep[INT/WFC Photometric $\\Halpha$ Survey of the Northern Galactic Plane,][]{Drew_Greimel.2005}, UVEX \\citep[The UV-Excess Survey of the Northern Galactic Plane,][]{Groot_Verbeek.2009} and the imminent VPHAS+ in the optical, which will collectively provide narrow and broadband photometry for the entire Galactic disc ($|b|<5$). Meanwhile, in the near infrared UKIDSS--GPS \\citep[UKIRT Infrared Deep Sky Survey -- Galactic Plane Survey,][]{Lucas_Hoare.2008} and VVV \\citep[Vista Variables in the Via Lactea,][]{Minniti_Lucas.2010} will, between them, cover much of the Galactic plane. The photometric surveys have been accompanied by spectroscopic surveys, including RAVE \\citep[][]{Siebert_Williams.2011} and SEGUE \\citep[][]{Yanny_Rockosi.2009}.  All these surveys have massively increased the volume of data available on our Galaxy and, with the right analytical tools, should allow us to develop an improved understanding of the Galaxy. This trend for the growth in survey astronomy is set to continue in the next decade; future facilities such as Gaia and LSST will gather data at a hitherto unseen rate. As such, an ongoing problem is how best to harness that data in order to extract as much information as possible \\citep[e.g.][]{Binney_only.2011}.  \\cite{Juric_Ivezic.2008} used stars from SDSS to study the stellar number density distribution at high Galactic latitudes (mostly $|b|>25^{\\circ}$, determining large scale properties of the Galaxy and discovering localised over densities in the Galactic stellar halo. \\cite{Ivezic_Sesar.2008} followed this by also analysing the metallicity of stars from SDSS. \\cite{Carollo_Beers.2010}, again using SDSS data, claim that they have discovered a second halo component. Though, as argued by \\cite{Schonrich_Asplund.2011}, it appears that this second component may be an artefact stemming from a substantial bias present in the distance estimates employed by \\cite{Carollo_Beers.2010}.  Studies, such as those discussed in the previous paragraph, demonstrate both the potential rewards and possible pitfalls associated with using large survey datasets. Similarly, it is becoming increasingly clear that simple methods of inference are potentially both ineffective and inaccurate. In large survey datasets various parameters influence data in a manner which is not directly obvious and different parameters will have conflicting effects on the observations. Consider the example of using star counts to estimate the scale length of the Galactic thin disc: a simple minded approach might consist of estimating the distance of the stars by `photometric parallax' and then measuring scale lengths from the resulting distribution. In this case, the individual distance estimates may well be biased and the sample will become incomplete and contaminated in a way which is not known. In short there are a number of confounding factors which prevent the scale length being accurately inferred through simple `inversion'.  Interstellar extinction is a particular hindrance when studying the Galaxy. It builds up gradually, along all lines of sight, in a manner which is not yet well known. A particular realisation of the difficulties interstellar extinction causes is the degeneracy between interstellar reddening and spectral type: an apparently red star may either be an intrinsically red late type star subject to little extinction or a heavily extinguished, intrinsically blue early-type star. Therefore, the presence of extinction makes it difficult to determine the spectral type and thus distance of stars, considerably complicating the study of structure in the Galaxy.  Furthermore, the 3D distribution of extinction is itself intrinsically interesting. Mapping extinction allows one to trace the distribution of dust in the Galaxy. Interstellar dust is itself only one component of the ISM, examining its distribution and comparing it to tracers of other components of the the ISM, such as CO maps \\citep[e.g.][]{Dame_Hartmann.2001} or HI maps \\citep[e.g.][]{Kalberla_Burton.2005}, offers a window onto the ISM and the physical processes which shape it.  \\cite{Juric_Ivezic.2008} were able to correct for the effects of extinction on their observed stars using the asymptotic Galactic extinction map of \\cite{Schlegel_Finkbeiner.1998}. This approach was reasonable as the majority of stars in their sample lie at high Galactic latitudes, beyond the Galactic dust layer, which is localised near the Galactic plane.  However, such an approach is not valid for the majority of stars in the Galaxy which lie close to the Galactic plane and are thus unlikely to be beyond all (or nearly all) the Galactic dust in their direction.  In light of the difficulties interstellar extinction causes and the intrinsic value of studying extinction, there has been an ongoing effort to map it over many years, going back to \\cite{Trumpler_only.1930} and \\cite{vandeKamp_only.1930} noticing the propensity for extinction to be stronger near the Galactic plane. Detailed discussions of the history of extinction mapping are available in \\cite{Sale_Drew.2009} and \\cite{Majewski_Zasowski.2011}. Recently, spurred on in part by the imminent launch of Gaia and construction of other telescopes, and the availability of high quality data, there has been a renewed focus on extinction mapping \\citep{Bailer-Jones_only.2011, Majewski_Zasowski.2011, Schlafly_Finkbeiner.2010, Schlafly_Finkbeiner.2011}.  Most recently \\cite{Berry_Ivezic.2011} have used SDSS and 2MASS photometry to estimate the spectral type, distance and extinction towards 73 million stars and so create extinction maps with fine angular scales. In common with many previous efforts to study extinction, they consider stars in isolation, fitting the parameters of a star without reference to the other stars local to it. In doing so they do not take advantage of the fact that properties of stars are correlated. For example, we know that stars located close to each other should be subject to similar extinctions: trivially we would be surprised if two stars closely located on the sky and apparently at similar distances exhibited the effects of vastly differing amounts of extinction. In not employing a method that utilises this information, their results are less precise than what could in principal be achieved with their data.  A contrasting approach is to use Galactic models to study extinction \\citep{Marshall_Robin.2006} and Galactic structure \\citep[e.g.][]{Robin_Creze.1992, Ruphy_Robin.1996, Sale_Drew.2010}, this offers several distinct advantages. Specifically, analysing stars en masse makes it possible to exploit the physical relationships that exist between stars. Moreover, by comparing Galactic models to real observations, effects such as Malmquist bias, sample contamination and sample incompleteness need not impact upon any inferences, as they will occur in a properly constructed Galactic model in a similar manner to how they occur in reality. However, such approaches typically involve binning the data in some manner. In doing so there is some inevitable loss of information, such that results obtained will be less precise than those obtained without binning. Additionally, such techniques marginalise over the parameters of individual stars, which may be undesirable if these are of interest.  \\subsection{Hierarchical Bayesian models}\\label{sec_intro_bayes}  Hierarchies pervade almost all aspects of Galactic astronomy. Trivially, the Galaxy is comprised of various stellar (e.g the thin and thick discs) and non stellar (e.g. dark matter halo, ISM) components. These can be broken down into further subcomponents, stars in the stellar case, such that the properties of the stars are dependant on the properties of the stellar component from which they are drawn. Other simple examples of hierarchies include: the distribution of the masses of stars in a newly formed star cluster which are drawn from some initial mass function (IMF) and the kinematics of stars in the Galaxy which are influenced by some global gravitational potential field.  Extinction mapping is a directly hierarchical problem: stars within a field trace a distance--extinction relationship for that sightline. However, existing methods of 3D extinction mapping do not exploit this hierarchical nature. For example, in methods such as those of \\cite{Neckel_Klare.1980} and \\cite{Arenou_Grenon.1992}, the distance and extinction of each star are first determined, without reference to other stars. Subsequently, the final distance-extinction relationship was estimated by fitting to the determined distance and extinction values for each star. Alternatively, a properly constructed method could use information gained from the field as a whole (i.e distance--extinction relationship) to help constrain the estimates of parameters of individual stars (i.e. their distance and extinction), which in turn can be used to refine our knowledge of the field. Using this form of `group knowledge', exploiting the correlations which exist between stars, enables parameters to be determined more precisely than is otherwise possible. Therefore, \\cite{Neckel_Klare.1980} and \\cite{Arenou_Grenon.1992}, in common with \\cite{Berry_Ivezic.2011}, obtain results less precise than their data are capable of.  Hierarchical Bayesian models are rich statistical models, they extend upon simple Bayesian models by allowing some parameters in the model to be dependant on other parameters. One can, in general, solve for all the parameters in the model, both those at a `lower' level, which may pertain to stars, as well as those at a `higher' level which could describe a sightline or the Galaxy.  On a purely mathematical level, if we have some observation $z$ and wish to estimate some parameter $\\theta$ we can employ Bayes' theorem: \\begin{gather} P(\\theta|z)=\\frac{P(z|\\theta)P(\\theta)}{P(z)}\\\\ P(\\theta|z) \\propto P(z|\\theta)P(\\theta) \\end{gather}  \\noindent It is also possible to define a hierarchical model, whereby $\\theta$ itself depends on a further parameter, often referred to as a hyperparameter, $\\phi$. Working from conditional probabilities one can obtain: \\begin{equation} P(\\theta, \\phi, z)=P(\\theta, \\phi|z)P(z) \\end{equation}  \\noindent Also: \\begin{align} P(\\theta, \\phi, z) &=P(z|\\theta, \\phi)P(\\theta, \\phi)\\\\ &=P(z|\\theta, \\phi)P(\\theta| \\phi) P(\\phi) \\end{align}  \\noindent Therefore: \\begin{equation} P(\\theta, \\phi|z) \\propto P(z|\\theta, \\phi)P(\\theta|\\phi)P(\\phi) \\end{equation}  \\noindent If $z$ is not directly conditional on $\\phi$, this then becomes: \\begin{equation} P(\\theta, \\phi|z) \\propto P(z|\\theta)P(\\theta|\\phi)P(\\phi) \\label{eqn:hierarchical_general} \\end{equation}  One could add a further tier to the hierarchy by introducing a further parameter on which $\\phi$ depends, this can be repeated as necessary. Also, it is possible to replace the single parameters or observations $z$, $\\theta$ and $\\phi$ with sets of several parameters.  It is instructive to consider a simple hierarchical model of two tiers: a top tier containing hyperparameters describing a galaxy and a lower tier of parameters describing the stars which inhabit the galaxy, of which we possess some observations. The form of hierarchical Bayesian model given by equation~\\ref{eqn:hierarchical_general} can be employed because any observations of stars we possess are not directly dependant on the state of the Galaxy. The power of the hierarchical model in this instance is clear, it allows information about the galaxy to be estimated directly from the data. That is to say, we can solve for the posterior distribution and from that we possess estimates of certain galaxy-wide parameters that would otherwise be difficult or impossible to determine. ",
        "conclusions": "\\label{sec_close}  Given the high frequency of hierarchies in Galactic astronomy, hierarchical Bayesian models will be a key tool for analysing the ever growing quantity of data on our Galaxy. The advantage of using hierarchical models is clear. Such models relate the parameters of many stars, through the characteristics of the Galaxy or the region in which they are found. Therefore, they provide additional information and allow the parameters which describe stars and those which describe the Galaxy to be determined more precisely and accurately. Conversely, methods which consider stars individually do not take advantage of the relationship between stars and the Galaxy they populate and cannot take full advantage of the data they employ.  Furthermore, employing a Bayesian methodology offers several advantages over classical frequentist techniques. Bayesian techniques avoid the classic inverse problem, whereby the parameters of interest are determined directly from the data, by instead finding the model which best describes the data. Such inversion can be particularly difficult when there are many parameters involved and their estimates are correlated. Additionally, with Bayesian techniques one is able to include prior information, in this case the basic physics of extinction and the accumulated historical knowledge of the Galaxy can both be included to give more precise and accurate results.  This paper has examined one potential use of hierarchical Bayesian models, employing them to map extinction in three dimensions. A good knowledge of where and how extinction builds up in the Galaxy is a major barrier to furthering our understanding of the structure an history of the Galaxy. There are few existing 3D extinction maps and those that do exist are limited by one or more of several factors, including the data they employ \\citep[e.g.][]{Neckel_Klare.1980} a reliance on a particular Galactic model \\citep{Marshall_Robin.2006}, by only being able to achieve a coarse distance resolution \\citep[e.g.][]{Majewski_Zasowski.2011} or by only treating individual stars in isolation \\citep[e.g.][]{Berry_Ivezic.2011}.  A particular algorithm, H-MEAD, which employs hierarchical Bayesian models has been described in detail. The model employed has been discussed, concentrating on the physical justification for its use, but also covering some of the practical aspects of its implementation. Subsequently, testing with synthetic photometry has demonstrated that H-MEAD does indeed produce results which are both accurate and precise. The precision of H-MEAD is massively improved with respect to MEAD \\citep{Sale_Drew.2009}, whilst the fact that H-MEAD is able to cope with features such as variation in stellar densities and metallicities will considerably reduce systematic errors.  It should also be remembered that H-MEAD not only maps extinction, but also provides precise estimates of the distances, masses, metallicities, etc of stars. There are many potential uses for such data, one possibility is studying the distribution of high mass stars in order to map the Galaxy's spiral structure.  In the longer term, it is not possible to consider the future usefulness of the method described in this paper and hierarchical models in general without mentioning future surveys. Gaia is expected to provide parallaxes for $\\sim10^9$ stars as well as spectroscopy for a subset. As such, it will clearly be central to any future analysis of the Galaxy. However, it will be far from the only data available. Rather than simply using Gaia data alone, a more precise analysis will be possible if data from existing and forthcoming surveys, be they photometric or spectroscopic, were used as well.  However, a major problem is this wealth of data should be analysed. Clearly, concealed within the data will be a gold-mine of information on our Galaxy and, in theory, it should prove possible to make spectacular advances in understanding the current and past state of the Galaxy. This massive advantage though becomes a major hurdle, in so far as the sheer quantity and complexity of the data will considerably complicate its analysis.  The method discussed in this paper, could be considered to be a prototypical example of how Gaia data might best be analysed. There are many potential improvements to the method, several of which have been discussed in section~\\ref{sec_extensions}, which will not only increase its accuracy but also extend the scope of what it can achieve. If applied this would make it possible to not only map extinction in 3D, but also map Galactic structure, estimate the Galaxy's star formation history and more besides."
    },
    "1208/1208.1614_arXiv.txt": {
        "abstract": "We use frequency-dependent position shifts of flat spectrum radio cores to estimate the kinetic power of AGN jets. We find a correlation between the derived jet powers and AGN narrow-line luminosity, consistent with the well-known relation for radio galaxies and steep spectrum quasars. This technique can be applied to intrinsically weak jets even at high redshift. ",
        "introduction": "Active Galactic Nuclei (AGN) are some of the brightest objects in the Universe. This makes them visible to huge distances, and therefore useful in applications ranging from studies of Galactic ionized gas (through interstellar scintillation; \\citet{LovellEA08}), to measuring positions on Earth to centimetre precision (via the technique of geodetic Very Long Baseline Interferometry; \\citet{MaEA09}), to cosmological applications (through their use as standard candles; \\citet{WatsonEA11}).  AGN play a key role in galaxy formation and evolution. The evolution of black holes and galaxies are tightly coupled throughout cosmic history \\citep{MagorrianEA98,HasingerEA05}, and nuclear activity of black holes imparts significant feedback on the surrounding gas through radiatively and mechanically driven outflows. Radio-emitting jets of relativistic particles inflate cocoons of radio plasma, which in turn drive bow shocks through the host galaxy and beyond \\citep{KA97}. On large scales, these shocks are observed to uplift the hot gas present in galaxy clusters. Even once jet activity ceases, jet--inflated radio lobes can rise buoyantly through cluster gas, transporting the gas outward \\citep[e.g.][]{FabianEA03,FormanEA05}.  AGN activity is intermittent, triggered by either radiatively efficient accretion onto the black hole (the so-called ``cold mode''); or radiatively inefficient (``hot mode'') accretion \\citet{BestHeckman12}. Powerful radio AGN are typically associated with the radiatively efficient mode, and often exhibit high-excitation ionization lines \\citep{HardcastleEA07}. This mode of AGN triggering is often associated with mergers \\citep{ShabalaEA12,RamosAlmeidaEA12}, and was likely prevalent in a younger, denser Universe. By contrast, the majority of low--redshift radio AGN are low power, and show no such strong emission lines; these are consistent with being fuelled by steady cooling of hot X--ray emitting gas in galactic haloes \\citep{BestEA05,PopeEA12}.  From a theoretical viewpoint, AGN feedback is a crucial ingredient of all galaxy formation models,  truncating (the otherwise) excessive star formation in massive galaxies at late times, and ensuring that present-day ellipticals are ``red and dead'' \\citep{CrotonEA06,BowerEA06}. This feedback can come from radiative pressure \\citep{FabianEA06}, quasar winds \\citep{NesvadbaEA08}, jets \\citep{SA09b} or a combination of these \\citep{HopkinsElvis10}.  A number of authors have argued that AGN feedback significantly affects the star formation histories of AGN hosts \\citep{KavirajEA11,AntonuccioDeloguSilk08,KauffmannEA03}. \\citet{RawlingsJarvis04} suggested on energetic grounds that a single AGN may be able to expel gas from multiple galactic haloes. Recently, \\citet{SKS11} confirmed this observationally: shocks from the largest (hundreds of kpc--scale) radio AGN were found to truncate star formation in dwarf galaxies that themselves may have never hosted an AGN. This potentially important new mode of feedback depends sensitively on the combination of AGN jet power and environment density. In order to form a large radio AGN, the radio jet needs to pierce the dense environment of the AGN host without being disrupted by magnetohydrodynamic instabilities. Weak jets in moderate density environments, or moderate power jets in dense environments, are easily disrupted within dense galactic cores \\citep{Alexander00}, and are thus unable to do large--scale feedback. Therefore, accurate measurement of AGN jet power is crucial to quantifying the effects of AGN on galaxy formation and evolution.  Measuring kinetic AGN power is difficult. Dynamical models, utilizing observed sizes and luminosities of radio AGN, can be used to estimate jet powers \\citep{KDA97,SAAR08,AntogniniEA12}. These models rely on some knowledge of environment into which the radio lobes are expanding, and AGN age. A closely related cavity power method relies on measurements of the work done by the radio lobes in inflating cavities in the X-ray emitting gas \\citep{RawlingsSaunders91,BirzanEA08}. This method also requires knowledge of both the density of the gas into which the lobes are expanding, and age of the radio lobes. Ages of synchrotron--emitting electrons can be estimated from either radio spectra \\citep{AlexanderLeahy87} or dynamical models, and the two estimates typically differ by a factor of two \\citep[e.g.][]{MachalskiEA04}. Furthermore, the cavity technique is limited to nearby, low-power radio galaxies, and the derived jet powers may thus be unrepresentative of the overall AGN population. Other sources of error in jet power estimates arise due to uncertainties in the exact gas density profile: for example, cavity measurements typically assume that the radio lobes expand into a constant density atmosphere throughout their lifetime.  In this paper, we apply a fundamentally different method for measuring AGN jet power to a sample of flat spectrum radio quasars. This method, first devised by \\citet{Lobanov98} and further discussed by \\citet{Hirotani05}, is based on the interpretation of the frequency dependent position shift of flat spectrum radio cores in terms of synchrotron self-absorption within the jet plasma. We compare the derived jet powers with narrow-line luminosities.  Throughout the paper, we adopt the concordance cosmology of $\\Omega_{\\rm M}=0.27$, $\\Omega_\\Lambda=0.73$, $h=0.71$. ",
        "conclusions": "We have calculated AGN jet powers from multi-frequency core shifts in flat spectrum quasars. We find a strong correlation between jet power and AGN narrow-line luminosity, consistent with the results of \\citet{RawlingsSaunders91} for steep spectrum radio galaxies. The core shift method has strong predictive power, allowing accurate estimates of jet power to be made in the absence of detailed jet and AGN orientation parameters.  The geometry of the Universe makes the core shift method a potentially very powerful one for measuring AGN jet powers, particularly in the era of new, high sensitivity instruments capable of very high astrometric precision such as the Square Kilometre Array \\citep{GodfreyEA12}. The angular diameter distance turns over at redshift 1.6, making it in principle even easier to measure a core shift at redshift 3 than at redshift 1.5. This is also true for other jet power measurement methods, such as dynamical and cavity models. However, the {\\it luminosity} distance increases with redshift, making the detection of steep spectrum objects difficult. The advantage of the core shift method is in the fact that the flat spectrum nature of core--shifting quasars allows Doppler boosting to bring even intrinsically faint radio sources above the detection limit. Thus, the core shift method will be able to probe higher redshifts and lower luminosities than traditional jet power measurement methods. In addition to measuring AGN jet power, this technique is potentially useful in studying the assembly and evolution of black holes in the very early Universe, and constraining the composition and the value of $\\gamma_{\\rm min}$ in parsec-scale jets."
    },
    "1208/1208.0338_arXiv.txt": {
        "abstract": "The {\\it Advanced Camera for Surveys}  (ACS) Fornax Cluster Survey is a Hubble Space Telescope program to image 43 early-type galaxies in the Fornax cluster, using the F475W and F850LP bandpasses of the ACS. We employ both 1D and 2D techniques to characterize the properties of the {\\it stellar nuclei} in these galaxies, defined as the central ``luminosity excesses'' relative to a S\\'ersic model fitted to the underlying host. We find $72\\pm13$\\% of our sample (31 galaxies) to be nucleated, with only three of the nuclei offset by more than 0\\farcs5 from their galaxy photocenter, and with the majority of nuclei having colors bluer than their hosts. The nuclei are observed to be larger, and brighter, than typical Fornax globular clusters, and to follow different structural scaling relations. A comparison of our results to those from the ACS Virgo Cluster Survey reveals striking similarities in the properties of the nuclei belonging to these different environments. We briefly review a variety of proposed formation models and conclude that, for the low-mass galaxies in our sample, the most important mechanism for nucleus growth is probably infall of star clusters through dynamical friction, while for higher mass galaxies, gas accretion triggered by mergers, accretions and tidal torques is likely to dominate, with the relative importance of these two processes varying smoothly as a function of galaxy mass. Some intermediate-mass galaxies in our sample show a complexity in their inner structure that may be the signature of ``hybrid nuclei'' that arose through parallel formation channels. ",
        "introduction": "\\label{sec:introduction}  Once viewed as relatively simple objects that formed in a single, ``monolithic\" collapse, early-type galaxies are now widely believed to have been assembled hierarchically through repeated mergers and accretions \\citep[e.g.,][]{white78,searle78,white91,kauffman00,cole00,sp05,bower06}. A property of most {\\it luminous} (e.g., $M_r \\lesssim -22.5$) early-type galaxies is that they appear to have formed the majority of their stars at high redshift ($z\\gtrsim1$, corresponding to ages of $\\tau \\gtrsim 7$-$8$~Gyr) and on short timescales ($\\Delta\\tau \\lesssim1$~Gyr) \\citep[e.g.,][]{bower92, franx93, thomas99b, trager00, wake06}. These features may be related to feedback from active galactic nuclei (AGN), which can generate jets and outflows that blow away gas and suppress star formation \\citep[e.g.,][]{silk98, king03, murray05, fabian06, robertson06}. The general trends in the star formation histories of low- and intermediate-luminosity early-type galaxies are not as well understood, but they are known to show considerable diversity and to depend sensitively on environment  \\citep[see, e.g.,][]{tolstoy09}.  The discovery of the $\\mathcal{M}_{\\rm BH}$--$\\sigma$ relation \\citep{ferrarese00,gebhardt00} points to a fundamental connection between the central black holes powering these AGN, and the dynamical properties of their host galaxies.  There are several other galaxy properties that have also been found to scale with black hole mass, including luminosity \\citep[e.g.,][]{kormendy95, ferrarese00}, light concentration \\citep[e.g.,][]{graham01}, global velocity dispersion \\citep[e.g.,][]{ferrarese00, gebhardt00, gultekin09}, bulge mass \\citep[e.g.,][]{magorrian98, marconi03, haring04}, and total gravitational mass of the host \\citep{bandara09}. Thus, it has become clear that an understanding of the central regions of galaxies, including black holes and AGN, is essential if we are to make sense of the formation and evolution of galaxies themselves.  However, the direct detection of black holes remains very challenging: see, e.g., Chapter~11 of \\citet{ferrarese05} for an overview of the observational difficulties. For kinematic measurements, a high central surface brightness is needed to obtain spectra of adequate S/N, and this requirement can pose problems for massive early-type galaxies with shallow surface brightness profiles in their cores. At the distances of the Virgo and Fornax clusters, the small angular size of the black hole ``sphere of influence\" in most galaxies introduces a further complication. For example, at 20~Mpc, the distance of Fornax, a black hole in a galaxy with $\\sigma=200$~km~s$^{-1}$ has a sphere of influence of only $0\\farcs2$ in radius (assuming the  $\\mathcal{M}$--$\\sigma$ relation from \\citealt{lauraa06}). It is therefore not surprising that  a dynamical black hole mass measurement exists for only a single early-type galaxy in the Fornax cluster \\citep[FCC~213;][]{houghton06, gebhardt07}.  \\defcitealias{cote06}{C06}   On the other hand, the correlation between a galaxy's mass and that of its black hole was recently shown to extend down to the central nuclear star clusters found in low-mass galaxies \\citep{laurab06, wehner06}. Other studies have reported similar relationships between black hole or nucleus mass and the host bulge luminosity, mass, and S\\'ersic index \\citep{rossa06, balcells07, graham07}. These results are suggestive of a global relationship between galaxies and {\\it both} types of central massive object (CMO; \\citealt[hereafter C06]{cote06}): however, it is still an open question as to whether black holes and nuclei form via the same mechanisms, or whether nuclei form first and serve as seeds for black hole formation.  The hydrodynamical simulations of \\citet{li07} of a shared formation mechanism for both nuclei and black holes via the gravitational collapse of gas in bulgeless disks were able to reproduce a CMO and host mass correlation even without imposing an  {\\em a priori} $M$-$\\sigma$ relation, and were observed to be in agreement with \\citet{laurab06}. Alternatively, \\citet{laurab06} noted that nuclei could, in principle, form in all galaxies, but in massive galaxies they might either collapse or be destroyed (or otherwise altered) by binary black holes. Using semi-analytic models, it was the demonstrated by \\citet{devecchi09} and \\citet{devecchi10} that nuclei could form at high redshifts and act as possible black hole seeds.  If nuclei and black holes form simultaneously, then it is possible that momentum feedback determines which object will eventually dominate the CMO mass. \\citet{mclaughlin06} noted that the same momentum flux that drives out gas from black holes \\citep{king03, king05} could also regulate the growth of nuclear star clusters. \\citet{nayakshin09} used this finding to explain why nuclei, not black holes, appear more likely to form in less massive hosts. Both objects can form simultaneously as gas is driven to the center of a galaxy through an event such as a merger, but it is the mass of the host bulge that sets the individual formation rates. Some evidence for such a scenario comes from observations of intermediate-luminosity galaxies \\citep{filippenko03, gonzalez08, seth08, graham09}, as well as a number of dwarfs \\citep{barth04, reines11}, that have been found to contain both a central stellar nucleus and a black hole. Indeed, using observations in the Virgo cluster, \\citet{gallo10} estimated that hybrid nuclei could occur in 0.3--7\\% of galaxies with stellar masses below $10^{11}M_\\odot$, and in less than 32\\% of hosts above this stellar mass. In short, the study of nuclei presents us with a new opportunity to deepen our understanding of how galaxies and black holes co-evolve.  Like black holes, nuclei pose some observational challenges of their own. Although their existence in some dwarf galaxies has been known for decades, comprehensive surveys of galaxy clusters --- in which the frequency of nucleation within complete galaxy samples could be robustly measured --- did not appear until \\citet{binggeli87} published their Virgo Cluster Catalog (VCC). This program observed 1277 members and 574 probable members of the Virgo cluster using the 2.5~m Las Campanas telescope;  about 26\\% of all dwarf galaxies in the VCC sample were found to be nucleated. Shortly thereafter, a similar survey of the Fornax cluster by \\citet{ferguson89} --- the Fornax Cluster Catalog (FCC) --- found nuclei in 103/249  $\\approx$ 41\\% of their dwarf galaxies. In the above studies, dwarf galaxies were identified primarily morphologically by their flat surface brightness profiles, although in general they were found to be fainter than $M_B\\simeq-18$~mag \\citep{sandage84}.  Given the low luminosities and small sizes of most of these nuclei, the frequencies of nucleation estimated from ground-based photographic studies are certainly lower limits. For instance, \\citet{lotz04} used WFPC2 on the Hubble Space Telescope (HST) to observe 69 dwarf elliptical galaxies in both Virgo and Fornax, finding nuclei in six galaxies that were previously classified as non-nucleated in the VCC and FCC.  Based on wide-field imaging of Virgo dwarfs from the Isaac Newton Telescope, \\citet{grant05} was able to identify many faint nuclei that were missed in the earlier photographic survey. In fact,  imaging of {\\it late-type} galaxies with HST commonly revealed ``nuclear clusters\" that had gone unnoticed in earlier studies, with an overall frequency of nucleation of $\\approx70\\%$ \\citep[e.g.,][]{carollo98, matthews99, boker04, walcher05, seth06}.   The first study to find a comparable frequency of nucleation among early-type galaxies was carried out by \\citetalias{cote06} with the Advanced Camera for Surveys (ACS) on HST: i.e., the ACS Virgo Cluster Survey \\citep[ACSVCS;][]{cote04}.\\footnote{Related papers from the ACSVCS on the central structure of early-type galaxies include \\citet{lauraa06, laurab06, cote07, glass11}.}  In addition to establishing a high frequency of nucleation for early-type galaxies (at least 66\\% for galaxies brighter than $M_B\\approx-15$), the high-resolution imaging made it possible to characterize the detailed properties of the nuclei for the first time, including their luminosity function, structural properties, color-magnitude relation, and nucleus-to-galaxy luminosity ratio. We note here that although in \\citetalias{cote06} and this work, we call the central excess of light rising above a galaxy's extrapolated outer surface brightness profile a ''nucleus``, these objects are not limited to being nuclear star clusters; certainly, some could be described as disks, bars, or other large scale structures, which have been observed by previous studies of early-type galaxy centers \\citep[e.g.][]{lauraa06, balcells07, morelli10}. In this paper, which is part of the ACS Fornax Cluster Survey (ACSFCS), we examine the properties of nuclei belonging to galaxies in the Fornax Cluster, which is located at a distance of $D = 20\\pm0.3\\pm1.4$~Mpc (statistical $+$ systematic error) \\citep{blake09}. This cluster is smaller, denser, more dynamically evolved, and more regular in shape than the Virgo cluster, and therefore allows us to study the properties of the nuclei of galaxies residing in a new and different environment.  \\defcitealias{jordan07a}{Paper~I} \\defcitealias{cote07}{Paper~II} \\defcitealias{glass11}{Paper~IV} \\defcitealias{blake09}{Paper~V} \\defcitealias{masters10}{Paper~VII} \\defcitealias{villegas10}{Paper~VIII} \\defcitealias{mieske10}{Paper~IX} \\defcitealias{liu11}{Paper~X}  Other papers in the ACSFCS series have described the data reduction procedures used in the survey \\citep[hereafter Paper~I]{jordan07a}, systematic variations in the central structure of galaxies \\citep[hereafter Paper~II]{cote07}, the logarithmic slope of the galaxy central surface brightness profiles \\citep[hereafter Paper~IV]{glass11}, and the use of surface brightness fluctuations as a distance indicator \\citep[hereafter Paper~V]{blake09}. Paper~III (2012, in prep.) of the ACSFCS will present a detailed isophotal analysis of the ACSFCS galaxies, including their dust properties, axial ratios, 2D structure, total magnitudes, colors, and surface brightness and color profiles. Papers studying the properties of globular clusters (GCs) in ACSFCS galaxies have examined their half-light radii \\citep[hereafter Paper~VII]{masters10}, luminosity function \\citep[hereafter Paper~VIII]{villegas10}, color-magnitude relation  \\citep[hereafter Paper~IX]{mieske10}, and color gradients \\citep[hereafter Paper~X]{liu11}.  The outline of this paper is as follows. In \\S\\ref{sec:observations} we describe the observations and methodologies used to measure photometric and structural parameters for the nuclei; in \\S\\ref{sec:results} we examine the nucleus properties, including their frequency of nucleation, luminosity function, sizes, surface brightness parameters, and colors; in \\S\\ref{sec:discussion} we put our results into the context of current formation scenarios; and in \\S\\ref{sec:summary} we summarize our main results. An appendix presents a comparison of 1D and 2D methods for measuring photometric and structural parameters of nuclei and their host galaxies. ",
        "conclusions": "\\label{sec:discussion}           \\subsection{The Role of Environment: Comparison to the ACSVCS} \\label{sec:compare}  As described in \\S\\ref{sec:introduction}, our Fornax survey was preceded by a similar study of 100 early-type galaxies in the Virgo cluster (ACSVCS, \\citealt{cote04}) where an  investigation into the properties of the nuclei in ACSVCS galaxies was carried out by \\citetalias{cote06}. Our prime motivation for a study of galaxies in the Fornax cluster is to provide a first glimpse into the properties of nuclei in two, rather different, clusters, and an assessment of the role played by environment in nucleus formation and evolution. The interested reader is referred to \\S~1 of \\citetalias{jordan07a}, which compares some key properties of the two clusters. Briefly, Virgo is overall a much larger cluster, with a mass almost 10 times that of Fornax (M$_{200} \\sim 4.2\\times10^{14}$~M$_\\sun$ vs. $(1\\mbox{--}7)\\times10^{13}$ \\citep{mclaughlin99, tonry00, drinkwater01}), and a velocity dispersion twice as large ($\\sigma_v\\sim760$ vs. $374$~km~s$^{-1}$ \\citep{binggeli87, drinkwater01}). Compared to the Virgo Cluster, Fornax is poorer (Richness Class 0 vs. 1, \\citep{abell89, girardi95}) and more compact (R$_{200} \\sim 0.7$ vs. $1.55$~Mpc). Its intracluster medium (ICM) has both lower temperature (1.20 vs. 2.58 keV) and metallicity (0.23 vs. 0.34 solar) \\citep{fukazawa98}, with the Fornax electron density at a given radius being about 1/4 that of Virgo \\citep{nulsen95, paolillo02}.  In this section, we will directly compare the results from both surveys. While \\citetalias{cote06} used King profiles for the nuclei in their paper, the ACSVCS results have since been updated with S\\'ersic model fits to the nuclei, which allows a fair comparison  between the two studies.\\footnote{See {\\tt https://www.astrosci.ca/users/VCSFCS/Data$\\_$Products.html}} Distances from \\citetalias{blake09} were used to calculate absolute magnitudes and physical sizes for both Fornax and Virgo galaxies. We note that the two surveys have slightly different cutoff magnitudes ($B_{T}\\simeq16$ for Virgo and $\\simeq15.5$ for Fornax) and that the distance modulus of Fornax is $\\sim$~0.5~mag larger than that of Virgo \\citepalias{blake09}. Therefore, the Virgo galaxies can reach absolute magnitudes roughly 1 mag fainter than those in Fornax.   \\begin{figure} \\figurenum{15} \\plotone{f15.eps} \\caption{Same as Figure~\\ref{fig:lum_hist}, but using absolute magnitudes and showing both ACSVCS and ACSFCS program galaxies (143 objects in total). \\emph{Top}: Luminosity distribution of the program galaxies for Virgo (solid and hatched red histograms) and Fornax (solid and hatched blue histograms). The solid histograms show the distribution of the 67 Virgo and 31 Fornax galaxies found to be nucleated by the ACSVCS and ACSFCS. \\emph{Bottom}: The percentage of galaxies found to be nucleated ($f_n$) for Virgo (red squares) and Fornax (blue circles).} \\label{fig:virgo1} \\end{figure}  \\subsubsection{Frequency of Nucleation}   In Figure~\\ref{fig:virgo1}, we plot the frequency of nucleation of the Virgo and Fornax program galaxies as a function of their absolute blue magnitude. The Virgo galaxies appear in red, and the Fornax galaxies are shown in blue. In the upper panel, we overlay histograms for all galaxies (hatched) and nucleated galaxies (solid). This figure demonstrates how the Virgo galaxy magnitudes extend to $\\sim1$ mag below those of Fornax, as explained above. Our Virgo sample contains 100 galaxies, 67 of which are found to be nucleated, so we obtain a total frequency of nucleation, $f_{n}=67\\pm8\\%$. This is in excellent agreement with the value of $f_{n}=72\\pm13\\%$ found for our full Fornax sample.  The bottom panel shows the frequency of nucleation in each luminosity bin. Both clusters exhibit very similar distributions with $f_{n}=0$ for the bright galaxies, while fainter than $M_{B}\\sim-19.5$, $f_{n}$ continuously stays above $\\sim70\\%$. Since our Virgo sample has 84 galaxies below $M_{B}=-19.5$, and our Fornax sample has 35,  we find the total frequency of nucleation for galaxies fainter than $M_{B}=-19.5$ to be $80\\pm10$\\% for Virgo and $89\\pm16$\\% for Fornax.  Both \\citetalias{cote06} and this study have shown that this sharp increase in frequency of nucleation compared to previous ground based studies (the VCC and FCC) is due mainly to surface brightness selection (see Figure~7 and 8 in \\citetalias{cote06} and Figure~\\ref{fig:galsb_vs_nucmag} in this work), which can be attributed to the improved resolution and depth offered by the ACS imaging. That is, the excellent angular resolution of  HST has allowed us to uncover previously undetected nuclei in both very high surface brightness galaxies, where the nuclei are difficult to distinguish from the main body, and low luminosity galaxies, in which the nuclei may lie below the magnitude limit of the older photographic surveys.  \\begin{figure} \\figurenum{16} \\centering \\leavevmode \\includegraphics[width = 0.85\\linewidth, trim =0cm 3.5cm 2cm 0cm, clip=true]{f16.eps} \\caption{Same as Figure~\\ref{fig:eta}, but using absolute magnitudes and including 67 ACSVCS and 31 ACSFCS nuclei. \\emph{Top}: Nucleus magnitude plotted against host galaxy magnitude, for the Virgo (red squares) and Fornax (blue circles) galaxies found to be nucleated, in the $g$- (\\emph{left}) and $z$-bands (\\emph{right}).  The lines show the weighted best fit relations, with the slope held fixed at unity (\\emph{solid}) and allowed to vary (\\emph{dashed}). The red and blue lines correspond to fits to the Virgo and Fornax samples respectively, while the black lines show the fits to the combined sample. \\emph{Bottom}: Nucleus-to-galaxy luminosity ratio, $\\eta$, against host galaxy magnitude, for the $g$-band (\\emph{left}) and $z$-band (\\emph{right}). The solid and dotted lines show the mean and one standard deviation, respectively, while the dashed line shows the best fit relation given by the dashed line in the upper panel, recast in terms of  $\\log(\\eta)$ and host magnitude.} \\label{fig:virgo2} \\end{figure}      \\subsubsection{Nucleus-to-Galaxy Luminosity Ratio} \\label{sec:virgoEta}  As  in \\S\\ref{sec:eta} and Figure~\\ref{fig:eta}, absolute nucleus magnitude has been plotted against absolute galaxy magnitude in the top panels of Figure~\\ref{fig:virgo2}. Relations of the form Equation~\\ref{eq:eta1} have been fitted using weighted least-squares to the Virgo and Fornax samples, both separately and combined, and the parameters and standard errors are recorded in Table~\\ref{tab:eta_virgo}, the results of which are in agreement to within the errors for both galaxy samples.  We also plot nucleus-to-galaxy luminosity ratio $\\eta$ as a function of absolute galaxy magnitude in the bottom panels of Figure~\\ref{fig:eta}. The values for the mean and standard deviation of $\\eta$ are given in Table~\\ref{tab:eta_virgo}. Taking the mean nucleus-to-galaxy luminosity ratio of both data sets combined, we obtain the following values for each band: \\begin{equation} \\begin{array}{rcl} \\langle \\eta_{g} \\rangle & = & 0.37 \\% \\pm 0.04 \\% \\\\ \\langle \\eta_{z} \\rangle & = & 0.34 \\% \\pm 0.04 \\%, \\\\ \\end{array} \\end{equation} which gives a mean value for both bands of \\begin{equation} \\langle \\eta \\rangle = 0.36 \\% \\pm 0.03 \\%. \\end{equation} The quoted errors refer the standard error on the mean.  Finally, we note that, due to the definition of $\\eta$, the best-fit relation from Equation~\\ref{eq:eta1} can be recast in terms of $\\log(\\eta)$ and galaxy magnitude, where $\\alpha_\\eta = -0.4 \\left(\\alpha_1 -1\\right)$ and $\\beta_\\eta = -0.4 \\beta_1$. This relation is plotted as the dashed line in the bottom panels of Figure~\\ref{fig:eta}, and we find that we do not see any significant trend between nucleus-to-galaxy luminosity ratio and galaxy magnitude.  \\begin{deluxetable}{lccrccccc} \\tabletypesize{\\footnotesize} \\tablecaption{Virgo and Fornax Nucleus-to-Galaxy Luminosity Ratios\\label{tab:eta_virgo}} \\tablewidth{0pt} \\tablehead { \\colhead{Sample} & \\colhead{Band} & \\colhead{$\\alpha_1$} & \\colhead{$\\beta_1$} & \\colhead{$\\beta_2$} & \\colhead{$\\langle\\log\\eta\\rangle$} & \\colhead{$\\sigma$} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{(mag)} & \\colhead{(mag)} & \\colhead{(dex)} & \\colhead{(dex)} } \\startdata ACSFCS   & $g$ & $0.90\\pm0.17$ & $3.99\\pm2.85$ & $5.78\\pm0.14$ & $-2.31$ & $0.32$ \\\\ ACSVCS   & $g$ & $0.81\\pm0.11$ & $2.79\\pm1.96$ & $6.12\\pm0.15$ & $-2.46$ & $0.47$ \\\\ Combined & $g$ & $0.80\\pm0.09$ & $2.68\\pm1.60$ & $6.04\\pm0.11$ & $-2.43$ & $0.44$ \\\\ ACSFCS   & $z$ & $1.06\\pm0.16$ & $7.38\\pm2.97$ & $6.22\\pm0.16$ & $-2.49$ & $0.35$ \\\\ ACSVCS   & $z$ & $1.02\\pm0.11$ & $6.51\\pm2.07$ & $6.21\\pm0.15$ & $-2.46$ & $0.53$ \\\\ Combined & $z$ & $1.02\\pm0.09$ & $6.63\\pm1.69$ & $6.21\\pm0.11$ & $-2.47$ & $0.49$ \\enddata \\end{deluxetable}  \\begin{figure} \\figurenum{17} \\plotone{f17.eps} \\caption{Same as Figure~\\ref{fig:lum_fun}, but using absolute magnitudes and including 67 ACSVCS and 31 ACSFCS nuclei. The luminosity functions for both the Virgo (red squares) and Fornax (blue circles) nuclei are shown, in the $g$-band (\\emph{top}) and $z$-band (\\emph{bottom}). Both data sets have been fitted with a normalized Gaussian.} \\label{fig:virgo3} \\end{figure}     \\begin{figure} \\figurenum{18} \\plotone{f18.eps} \\caption{Same as Figure~\\ref{fig:size_hist}, but using parsecs and including 67 ACSVCS and 31 ACSFCS nuclei. The distribution of half-light radii for both Virgo (red squares) and Fornax (blue circles) nuclei are shown. The red and blue vertical dotted lines indicate the adopted resolution limit of $\\sim0\\farcs025$, which corresponds to 2.0~pc in the ACSVCS and 2.4~pc in the ACSFCS.}\t\\label{fig:virgo4} \\end{figure}   \\begin{deluxetable}{lccc} \\tabletypesize{\\small} \\tablecaption{Virgo and Fornax Nucleus Luminosity Function\\label{tab:lum_fun_virgo}} \\tablewidth{0pt} \\tablehead { \\colhead{Sample} & \\colhead{Bandpass} & \\colhead{$\\bar{m}^0_n$} & \\colhead{$\\sigma_n$} \\\\ \\colhead{} & \\colhead{} & \\colhead{(mag)} & \\colhead{(mag)} } \\startdata ACSFCS & $g$ & $-11.54\\pm0.03$ & $1.58\\pm0.02$ \\\\ ACSVCS & $g$ & $-11.45\\pm0.02$ & $1.87\\pm0.02$ \\\\ ACSFCS & $z$ & $-12.67\\pm0.03$ & $1.70\\pm0.02$ \\\\ ACSVCS & $z$ & $-12.80\\pm0.02$ & $2.16\\pm0.02$ \\enddata \\end{deluxetable}  \\subsubsection{Nucleus Luminosities and Sizes}  In Figure~\\ref{fig:virgo3}, histograms of nuclei luminosities for both our Virgo and Fornax sample are compared. The parameters of the weighted maximum-likelihood fit of a normalized Gaussian to each sample are given in Table~\\ref{tab:lum_fun_virgo}, where the errors on the fitted parameters are the standard errors. Although we find differences between $\\bar{m}^0_n$ between the two surveys, there amounts are comparable to the errors estimated for the nuclei magnitudes.  We plot a histogram of nucleus sizes in Figure~\\ref{fig:virgo4} for both our Virgo and Fornax samples. Although there is a large range in size (the very large Virgo nucleus belongs to VCC~1178), most nuclei appear to have radii $<10$ pc. The typical sizes are in good agreement, with median values of $5.7$~pc in the $g$-band for both clusters, and $7.2$~pc and $7.0$~pc in the $z$-band for Virgo and Fornax respectively.  \\subsubsection{Other Properties}  In addition to the above properties, we find the Virgo and Fornax nuclei to be remarkably similar in a number of other ways. First,  and most obviously, both the ACSFCS and ACSVCS  galaxies  exhibit a trend along the luminosity function in which their central surface brightness profiles gradually change from having a luminosity ``deficit'' to an ``excess'': see, e.g., Figures~3 and 4 in \\citetalias{cote06}, Figure~1 of \\citetalias{cote07}, Figures~\\ref{fig:acs_images} and \\ref{fig:sb_profiles} here, as well as a detailed discussion of this trend in \\citetalias{glass11}.   Plotting surface brightness against magnitude, the nuclei are found to have different scaling relations than the GCs (see Figure~18 in \\citetalias{cote06} and Figure~\\ref{fig:scaling} here). Although \\citetalias{cote06} used integrated nucleus colors in their study, our use of aperture colors may be a more appropriate comparison to the King profiles used to determine the integrated nucleus magnitudes. Indeed, the best-fit line parameters outlining the color-magnitude relation for the nuclei with $B_{T}\\leq13.5$, given by Equation~13 in \\citetalias{cote06} and Equation~\\ref{eq:col} in this work, are in good agreement.   Overall, we find a striking similarity between the nuclei of Virgo and Fornax, despite the clear environmental differences between the two clusters. This agreement suggests that the physical characteristics of individual galaxy clusters (such as ICM density), or the processes that depend on them (such as ram pressure stripping efficiency), do not play a dominant role in the formation and evolution of nuclei in early-type galaxies. Thus, it seems we can consider the nuclei examined here as being representative of those in early-type galaxies in general.   \\subsection{Extension to Low Luminosity: Comparison to the Local Group}  Although the ACSVCS and ACSFCS provide a reliable measurement of the nucleation frequency for galaxies brighter than $M_{B}\\lesssim-15$~mag, it is instructive to consider the frequency of nucleation in galaxies fainter than this magnitude limit. We can do so by examining the members of the Local Group, where the smallest observed dwarf galaxies reach magnitudes faint as $M_{V}=-1.5$~mag and can have effective radii on the order of $\\sim30$~pc \\citep[see, e.g.,][]{martin08}. As sample completeness is a concern for such faint, compact systems, we focus on the subset of early-type galaxies brighter than $M_B \\approx -8$.  At present, there are 25 known early-type galaxies in the Local Group brighter than this limit \\citep[compiled from][]{mateo98, mcconnachie05, brasseur11}. Of these, only two (NGC~205 and M32) are brighter than the ACSVCS limiting magnitude of $M_{B}=-15$~mag \\citep{mateo98}, both of which are known to be nucleated \\citep[e.g.][]{kent87, lauer98, mateo98, butler05, derijcke06}. Moving down the luminosity function, at most six other galaxies may contain either nuclei or kinematically/structurally distinct features near their core, listed in order of decreasing luminosity: NGC 147 \\citep{derijcke06}, Sagittarius \\citep{mateo98, layden00, monaco05, bellazzini08}, Fornax \\citep{coleman04, coleman05, coleman08}, Sextans \\citep{kleyna04, walker06}, Andromeda~II \\citep{mcconnachie06}, and Ursa~Minor \\citep{kleyna03, palma03}.\\footnote{Although it is traditionally classified as non-nucleated, we include the Fornax dSph in this list since GC \\#4 is located $\\sim$ half a core radius from the galaxy photocenter \\citep[see Figure~1 of][]{coleman08} and might thus be classified as a dwarf with an offset nucleus if moved to the distance of the Virgo or Fornax clusters.} Considerable caution is advisable here since, in some cases (e.g., in Ursa Minor and, especially, in Sextans), the ``nuclei\"  are rather subtle substructures (sometimes only apparent with the addition of kinematic data) that bear little resemblance to the prominent, compact nuclei seen in the faintest ACSVCS and ACSFCS galaxies. Yet, even with this liberal definition of a ``nucleus\", only 8 out of the 27 Local Group early-type galaxies ($f_n = 30\\%$) can be classified as nucleated. If we exclude Fornax, Ursa Minor and Sextans from the list of nucleated galaxies, then $f_n$ falls to 19\\%. While it is possible that some nuclei have yet to be discovered, it seems certain that many of the faint Local Group galaxies do {\\it not} contain a nucleus; e.g., \\citet{mcconnachie06} studied of structural properties of six Andromeda satellites using deep, homogeneous imaging, and found a nucleus in only a single object (And~II).  We conclude that the frequency of nucleation along the Local Group sample is clearly far lower than in either our ACS surveys of the Fornax or Virgo clusters. Why is there such a large disparity in $f_n$? We speculate that the lack of nuclei in very faint galaxies could be related to the general absence of GCs in galaxies below $M_B \\sim -12$ \\citep[see, e.g.,][]{peng08}. If nuclei in low-mass galaxies are assembled through GC infall and mergers (see \\S\\ref{sec:gcinfall}), then the presence of GCs would obviously be a prerequisite for nucleus formation. The faintest galaxies in the Local Group known to contain GCs are Sagittarius   and Fornax, with $M_{B}= -12.8$ and $-12.6$ respectively \\citep{mateo98}. The former is unquestionably nucleated, while Fornax {\\it may} meet the definition of a nucleated galaxy (see above). Because no Local Group dwarfs below this magnitude are known to contain GCs, such galaxies might have been unable to form a nucleus if star cluster infall is the dominant mode of nucleus formation in low-mass systems.  It is also interesting to note that, assuming a constant nucleus-to-galaxy luminosity ratio of 0.4\\%, then the expected nucleus magnitude of a $M_B = -12.6$~mag host would be $M_B=-6.6$~mag. This corresponds closely to the mean turnover magnitude of the globular cluster luminosity function, $M_V\\approx-7.5$~mag \\citep[e.g.,][]{jacoby92, harris01, brodie06}, suggesting that galaxies may be unable to form nuclei at the point where the expected nucleus luminosity would fall below the typical GC luminosity.  However, as caveats we firstly note that the nucleus of Sagittarius \\citep{monaco09} as well as the very central region of the Andromeda satellite NGC~205 \\citep{siegel07} have been observed to have undergone multiple star formation episodes, which indicates that other processes in addition to GC accretion must have shaped their formation history. In addition, the nuclei late-type dwarfs  have been shown to \\emph{not} form form exclusively from GC infall \\citep[e.g.,]{walcher06} or gas accretion \\citep{hartmann11}, even though it has been observed that GC specific frequency is independent of morphology \\citep{georgiev10} and thus should be the same for both early- and late-type dwarfs.   \\subsection{Formation and Evolution Models}\\label{sec:models}  The origin of nuclei remains an open theoretical problem, with two main avenues of nucleus formation presently considered most viable. The first proposes that a galaxy's star clusters will experience orbital decay due to dynamical friction and spiral inwards, eventually coalescing at the center of the galaxy. The second formation mode focuses on gas accretion at the center of the galaxy, followed by star formation. Some similarities in the scaling relations of nuclei and black holes  (see \\S\\ref{sec:introduction}) have also given rise to models that consider the formation of both types objects in a shared context. In this section, we shall examine theoretical studies of nuclei formation in light of our new results, as well as models that explore the relationship between nuclei and black holes.   \\subsubsection{Dissipationless Infall of Star Clusters} \\label{sec:gcinfall}  \\citet{tremaine75} first suggested that the nucleus of M31 was formed from GCs that spiraled inward to the galaxy center due to dynamical friction, and this mechanism continues to offer an attractive explanation for the assembly of nuclei in at least some galaxies. Of course, not all clusters that come close to the center of a galaxy will necessarily contribute to the formation, or growth, of a stellar nucleus; as \\citet{capuzzo93} showed, dynamical friction and tidal stripping are competitive processes, where GCs are more readily destroyed by large nuclei, limiting nucleus growth.  Nevertheless, some fraction of GCs are expected to avoid tidal disruption and could contribute to either nucleus formation, or the growth of pre-existing nuclei. Evidence in favor of this process was described in \\citet{capuzzo99}, who pointed out that the radial distribution of GCs in galaxies is less centrally concentrated than the halo stars (see also \\citealt{mclaughlin95, mclaughlin99, cote01, cote03, peng08}). Such ``missing'' clusters could have contributed to nucleus formation. Monte Carlo simulations based on this premise by \\citet{lotz01} predicted nuclei luminosities for dEs with $-17\\lesssim M_B \\lesssim -12$ that were consistent with observations for the brighter galaxies within this range, although they were overestimated for less luminous ones. The over-prediction of nuclear luminosities in their low-mass systems resulted from their short dynamical times --- meaning that nuclei are able to grow very efficiently --- in spite of the fact that these galaxies have relatively few star clusters (see e.g., \\citealt{peng08}).  Numerical simulations by \\citet{oh00} and similar, higher resolution N-body simulations by \\citet{capuzzo08a, capuzzo08b} were able to successfully reproduce the observed surface brightness profiles of known nucleated galaxies. A dependence on local tidal field was found in the \\citet{oh00} model, where disruptive tidal forces on the outskirts of galaxy clusters would alter GC orbits, increasing dynamical friction timescales and decreasing nucleation frequency. The \\citet{capuzzo08a, capuzzo08b} models suggest that, if linear scaling is assumed, then the observed nuclei could have formed from the infall of tens, to hundreds, of GCs (see also \\S4.9 and \\S5.2.4 of \\citetalias{cote06}). Both simulations found that nuclei may begin to coalesce away from the galaxy photocenter, although to quite different extents: i.e., up to $\\sim0.3$~kpc and settling within $\\sim1$ Gyr in \\citet{oh00}, and $\\sim4$ pc away in \\citet{capuzzo08a, capuzzo08b}.  Other simulations by \\citet{bekki04} observed that the scaling relations of nuclei formed through mergers of GCs would be notably different than those of the GCs.  In \\S\\ref{sec:scaling}, we discussed that the predicted scaling relation for nuclei in these simulations, $R_e \\propto {\\cal M}_*^{0.38}$, was generally in good agreement with observations (see Figure~\\ref{fig:scaling}). More recent work by \\citet{bekki10a} focused on simulations of star cluster infall due to dynamical friction in {\\it disk} galaxies. He found that the effectiveness of dynamical friction did not depend strongly on bulge mass, but increased with smaller disk mass, and with larger disk mass fraction, galaxy surface brightness, and star cluster mass. The ratio of nucleus mass to disk mass was found to decrease as a function of increasing disk mass, with a mass ratio of $\\gtrsim0.4\\%$ for smaller disks, and $\\lesssim0.1\\%$  for disks with masses $M\\gtrsim10^{9}M_{\\sun}$. However, star cluster mergers on to a disk may not be sufficient to explain nuclei formed in $M_V \\sim -19.5$ spirals. N-body simulations by \\citet{hartmann11}, which aimed to reproduce the observed kinematics of the nuclei in M33 and NGC~4244, found that star cluster accretion on to a disk did not produce the necessary line-of-sight velocity rise, and at least half of the nucleus mass had to come from gas dissipation.    \\begin{figure} \\figurenum{19} \\plotone{f19.eps} \\caption{\\emph{Top}: Histogram of masses for the 143 galaxies from the ACSVCS and ACSFCS surveys. \\emph{Bottom}: Dynamical friction timescales, T$_{\\rm DF}$, plotted as a function of galaxy mass. Two sets of curves are shown. The dashed blue curves show calculations for initial GC radii, $R_i$, equal to the galaxy effective radii (see the lower panel of Figure~\\ref{fig:scaling}), while the dotted red curve shows $R_i$ fixed to 1.3~kpc, the median effective radius for ACSFCS galaxies. In both cases, T$_{\\rm DF}$ is plotted for five GC masses: $0.1, 0.25, 0.5, 1.0$ and $2.0$ million solar masses. Note the sharp decline in T$_{\\rm DF}$ for low-mass galaxies.} \\label{fig:dyn} \\end{figure}   Some provisional evidence for dissipationless formation in at least some galaxies was presented in \\citet{paudel2011}, who used optical spectroscopy for Virgo cluster dwarfs to study both their stellar populations and those of their nuclei. Despite the small sample and the different environment (Virgo vs. Fornax), their data present an interesting opportunity to speculate on possible formation mechanisms for the ACSFCS nuclei. \\citet{paudel2011} found that nuclei in a handful (5) of the faint ($-16 \\lesssim M_B \\lesssim -14$) galaxies in their sample were older and more metal poor than their hosts, which is certainly suggestive of a connection to GCs. At higher luminosities, most of their nuclei were found to be {\\it younger} than their hosts.  While inconsistent with nucleus formation from old globular clusters, this observation may still be compatible with cluster infall, as our observations and many others have shown that ongoing star cluster formation can be present throughout some galaxies \\citep[e.g.,][]{anders04, kyeong10}. In the ACSFCS sample, FCC~119, FCC~90 and FCC~26 are possible examples of $M_B > -19.5$ galaxies with young cluster systems.   Additional support for such a scenario may come from the GC luminosity functions in Virgo and Fornax galaxies. The widths of GC luminosity functions are known to decrease significantly with galaxy luminosity, a trend that is accompanied by a slight decrease in turnover mass  (\\citealt{jordan06, jordan07b}; \\citetalias{villegas10}). This truncation of the GC population on the bright end of the luminosity function may be caused, at least in part, by the shorter dynamical friction times as galaxies become less massive, although other (external) processes could also play a role (see \\S7.2 of \\citealt{jordan07b}).  We revisit the question of star cluster infall efficiency by calculating the dynamical friction timescale, T$_{\\rm DF}$, for all galaxies in our ACS surveys of Fornax and Virgo. The upper panel of Figure~\\ref{fig:dyn} shows the distribution of galaxy masses from the combined surveys (filled histogram), while the lower panel show the dependence of T$_{\\rm DF}$ on galaxy mass, ${\\cal M}_*$, which is given by $${\\rm T}_{\\rm DF}  =  {2.64\\times10^2 \\over {\\ln{\\Lambda}}}\\biggl({R_i \\over 2~{\\rm kpc}}\\biggr)^2\\biggl({v_c \\over 250~{\\rm km~s^{-1}}}\\biggr)\\biggl({10^6{\\cal M}_{\\odot} \\over {\\cal M_{\\rm GC}}}\\biggr)~{\\rm Gyr}.\\eqno{(16)}$$ Here $R_i$ is the initial galactocentric radius of the star cluster, $v_c$ is the circular velocity of the (assumed isothermal) galaxy, and ${\\cal M}_{\\rm GC}$ is the mass of the star cluster \\citep{binney08}. In this equation, $\\ln{\\Lambda}$ is the coulomb logarithm, which is defined as $${\\ln{\\Lambda}}   =   \\ln{\\biggl[{b_{\\rm max}v_c^2 \\over G({\\cal M_{\\rm GC}} + \\textsc{m})}\\biggr]} \\eqno{(17)}$$ where $b_{\\rm max}$ is the maximum impact parameter between the cluster and the interacting particle (a star of mass $\\textsc{m}$). Following \\citet{lotz01}, we assume $v_c \\simeq \\sqrt{2}\\sigma$ where $\\sigma$ is the integrated-light velocity dispersion measured within $R_e/4$ from McLaughlin et~al. (2012, in prep.). We also take $b_{\\rm max} = R_e$ for all galaxies, with $R_e$ measured directly from the ACS imaging (see \\S\\ref{sec:scaling} and Figure~\\ref{fig:scaling}).  Calculations have been carried out for five different star cluster masses (i.e., 0.1, 0.25, 0.5, 1 and 2 million solar masses)\\footnote{Recall that in the Milky Way, the GC mass corresponding to the peak of the luminosity function is $2.4\\times10^5{\\cal M}_{\\odot}$ \\citep{mclaughlin99}.} and for two assumptions for $R_i$. In the first case, we take $R_i = R_e$ (see also \\citealt{lotz01}) which is shown as the dashed blue curves in Figure~\\ref{fig:dyn}. In the second case, we simply fix $R_i$ at  the median effective radius, $1.3$~kpc, for all galaxies in the ACSFCS sample. The results in this case are indicated by the dotted red curves in Figure~\\ref{fig:dyn}. Although T$_{\\rm DF }$ clearly varies with the assumed cluster mass and the precise choice of $R_i$, the strong mass dependence noted by previous investigators is clearly apparent in this figure. In particular, the dynamical friction timescales are dramatically shorter in galaxies with ${\\cal M}_* \\lesssim10^{10}{\\cal M}_{\\odot}$ compared to higher-mass galaxies. We conclude that star cluster infall seems like a viable, indeed a likely, candidate for the growth of nuclei in low- and intermediate-mass galaxies in our sample. For the highest-mass galaxies, the mechanism appears much less viable given the fact that, in these systems, T$_{\\rm DF}$ greatly exceeds the Hubble Time for all but the most massive and centrally concentrated star clusters.   Finally, we conclude this section with some final remarks on Figure~\\ref{fig:scaling}, which compared the structural parameters of nuclei to those of  GCs and their host galaxies. While there is, as noted in \\S\\ref{sec:scaling}, good agreement with the nuclei size-mass relationship found by \\citet{bekki04} from simulations of GC mergers, there are reasons to believe that a single relation cannot be appropriate for all nuclei which, in our sample, span more than four decades in mass.  For comparison, the simulated nuclei of \\citet{bekki04} span a factor of just ten in mass. It is to be expected that the precise form of the size-mass relation in the context of the GC merger model will be different in different mass regimes. For instance, when only a small number of mergers contribute to the nucleus, we expect from the virial theorem and conservation of energy that $R_e \\propto {\\cal M}_*^{0.5}$. At later times, when the mass of the nucleus greatly exceeds the mass of an accreted GC, the relation should steepen to $R_e \\propto {\\cal M}_*$. These scaling relations, shown in the lower panel of Figure~\\ref{fig:scaling}, are in good agreement with the observed sizes and masses.  All in all, based on the existing data, we believe that cluster infall must have played an important role in the formation of the nuclei the low- and intermediate-mass hosts within our sample. At the same time, the red colors of some of the largest and most massive nuclei (\\S\\ref{sec:colors}) present a strong challenge to this model, suggesting that an additional process --- most likely the dissipational infall of metal-rich gas  --- likely begins to dominate the formation of nuclei in galaxies of progressively larger masses (\\citealt{mihos94}; \\citetalias{cote06, cote07}; \\citealt{hopkins08, hopkins09a}).   \\subsubsection{Dissipational Infall of Gas} \\label{sec:gas}  It has long been suspected that nuclei could form through star formation following the accretion of gas in galaxy centers \\citep{vandenbergh86}, although the exact origin of the gas, and the mechanism that triggers the inflow, are matters of debate.  In some models, the gas is assumed to originate from outside the galaxy. \\citet{davies88} proposed that dEs may be formed from fading stellar populations in dwarf irregulars, where the accretion of $\\textsc{Hi}$ gas induced starbursts, the final one occurring in the center and forming the nucleus. \\citet{silk87} predicted that the intergalactic medium (IGM) could fall into dwarf galaxies when it is cooled and compressed during group formation. This model noted that dwarfs closer to large galaxies may not be able to form nuclei as efficiently, since the large galaxy's tidal field makes it difficult for the dwarf to capture the gas. \\citet{babul92} found an opposite trend with environment: they observe that nucleus evolution may depend on local IGM density, because this determines whether supernova-driven gas outflows are able to escape. Dwarfs in low-pressure regions would have their gas ejected and then fade away, while winds in dense environments would be restricted to the starburst region by the IGM. This confinement could cause gas to cool and recollapse, creating two short or one prolonged starburst.  Gas might also be funneled to the centers of galaxies which have disks and axisymmetric features.  \\citet{milosavljevic04} suggested that in spiral galaxies, magneto-rotational instability in the disk transports gas to the center. \\citet{bekki06} and \\citet{bekki07} performed chemodynamic simulations of the inner $1$~kpc of dwarf galaxies with stellar masses  of $2.5\\time10^7 \\leq M_{\\rm sph} \\leq 1.0\\times10^9$, to explore the remnant created through dissipative merging of stellar and gaseous clumps formed from nuclear gaseous spiral arms in a gas disk. The simulations produced nuclei which that rotating and flattened, consisting of stars with varying ages and metallicities. Although the initial clump was found to form off-center (about $200$~pc by visual inspection of the simulation data), it would fall into the center within $100$~Myr. They found that overall, the nuclei were characteristically younger and more metal rich than the host, with more massive hosts creating more metal-rich nuclei. Gas settling timescales increased with decreasing dwarf mass (due to feedback being more effective in smaller galaxies), so low mass dwarfs were found to have younger and bluer nuclei. More massive and dense nuclei were formed in more massive dwarfs with deeper central potentials, and both the mass and mass fraction of the nucleus were found to increase with spheroid mass. Nuclei in high surface brightness galaxies should also have higher surface brightness, owing to the increased dynamical friction due to higher stellar densities. The nucleus  surface brightness was strongly dependent on the gas fraction of the host, and  thus may be more likely to form in this manner in  late-type galaxies with relatively large amounts of gas. Finally, the addition of a central black hole to the simulation had little effect on the properties of the remnant nucleus.  Another source of nuclear material, which was first proposed by \\citet{bailey80}, could arise from stellar winds. It was found that only a small ($\\sim10^6 M_{\\sun}$) amount of gas was needed to cause an inflow for an elliptical galaxy with $M_{\\rm gal}\\sim10^{11} M_{\\sun}$. \\citet{seth10b} observed that such a mechanism could produce the age, abundance gradient, and rotation curve seen in the nucleus of M32.  The dissipative infall of gas to the galaxy center can also be induced by galaxy mergers. \\citet{mihos94} performed N-body simulations of disk galaxy mergers, where they found that gas dissipation and the star formation that followed created dense stellar cores in the remnant.  Similar higher resolution simulations by \\citet{hopkins08, hopkins09a}, showed that gravitational torques during gas-rich mergers removed the angular momentum of the gas, which would then undergo gravitational collapse. The amount of gas infall was found to largely depend on the progenitor galaxy gas fraction, while the addition of a central black hole was not found to have a significant effect on the the properties of the final remnant. Unfortunately, these models lacked the resolution to study typical nuclei, particularly those in the low-mass galaxies: i.e., apart from a small number of cE galaxies in the ACSVCS sample, which have likely been heavily tidally stripped (e.g., \\citealt{faber73, lauraa06, cote08, chilingarian09b, huxor11}; McLaughlin et~al. 2012, in prep.), the simulated galaxies of \\citep{hopkins09a} have masses $\\gtrsim10^{10}{\\cal M}_{\\odot}$, more than ten times larger than the masses of the faintest galaxies in the ACS surveys. However, in this restricted mass regime, the properties of these simulated galaxies are in good agreement with our ACSFCS (and ACSVCS) observations.  Likewise, the simulations of \\citet{bekki06} and \\citet{bekki07}, which instead focused on the {\\it low-mass} galaxies, also appear to be consistent with observations, including those from our HST/ACS imaging and results from ground-based spectroscopy. First, the nuclei in these simulations were found to be younger and more metal rich than their hosts, with nucleus metallicity increasing with host mass, a trend that was seen in \\citet{paudel2011}. Second, their finding that low-mass dwarfs have younger and bluer nuclei is consistent with some of the nuclei from \\citet{paudel2011}, as well as with the nucleus colors observed in our study.  Finally, they also found that the mass fraction of the nucleus increased with host spheroid mass, and that their simulated surface brightness profiles showed nuclei which become more prominent with increasing dwarf mass, whereas in low-mass dwarfs the nuclei were barely distinguishable. It is therefore possible, as discussed in \\S\\ref{sec:gcinfall}, that nucleus formation through gas infall may be most significant for intermediate- and high-mass galaxies. In their analysis of the ACSVCS, \\citetalias{cote06} noted that some of the reddest and brightest nuclei ``may be candidates for the {\\it dense stellar cores} that form in numerical simulations \\citep{mihos94} when (chemically enriched) gas is driven inward, perhaps as a result of mergers.\" Such a result can be reconciled with our nearly constant nucleus-to-galaxy luminosity ratio if star cluster infall accounts mainly for nucleus build up in lower-mass galaxies. At intermediate masses, both processes could contribute significantly to the growth of nuclei; candidates for such {\\it hybrid nuclei} in the ACSFCS include FCC~43, FCC~249, FCC~310, FCC~148 and FCC~301, which may consist of both compact and extended components.  This basic scenario is also consistent with the general view that mergers (which can drive gas to the central regions of a galaxy) become increasingly important as galaxy luminosity increases, a consequence of the hierarchical merging paradigm. The observation that galaxy concentration --- parameterized by S\\'ersic index $n$ --- varies smoothly with galaxy luminosity (e.g., \\citealt{jerjen97, graham03a, lauraa06}; McLaughlin et~al. 2012, in prep.; see also \\S\\ref{sec:sbp}) provides strong supporting evidence for this picture, as violent relaxation of merger remnants is thought to be responsible for the creation of de Vaucouleurs profiles \\citep[e.g.][]{barnes88,barnes92}, while S\\'ersic index of both bulge and disks of spirals has been shown to increase after satellite infall \\citep{eliche05}. Figure~\\ref{fig:sersicindex} shows the dependence of two fundamental parameters for nuclei --- luminosity fraction and effective radius --- against host galaxy S\\'ersic index (Paper~III). Those galaxies whose internal structure has likely been transformed most extensively through mergers, accretions and harassment (i.e., those galaxies with high S\\'ersic indices) tend to have the most luminous and spatially extended nuclei (although the trend between $n$ and $\\eta$ is statistically significant only when an unweighted fit is used). These trends are generally consistent with an increasing importance for gas dissipation as ones moves to higher and higher mass galaxies. Nuclei formed through merger-driven gas inflow could also be expected to follow a mass-radius scaling relation, as \\citet{hopkins10b} found that stellar systems may have a maximum stellar surface density, due to feedback from massive stars.  \\begin{figure} \\figurenum{20} \\plotone{f20.eps} \\caption{\\emph{Top}: Nuclei luminosity fraction plotted against S\\'ersic index of the host galaxy, $n_{\\rm gal}$. \\emph{Bottom}: Nuclei effective radius as a function of $n_{\\rm gal}$. The dashed line in each panel shows the weighted best-fit linear relation; unweighted fits are shown by the dotted lines. The nuclei in both clusters show weak trends with S\\'ersic index (or, equivalently, galaxy mass) in the sense that the central ``excess\" above the fitted S\\'ersic model seem to be brightest and largest in galaxies with the largest $n_{\\rm gal}$. These galaxies have likely undergone fewer mergers and accretions than those with $n_{\\rm gal} \\sim 1$.} \\label{fig:sersicindex} \\end{figure}   One complication with the gas inflow model is that it obviously requires the presence of gas, which is not consistent with the ``classical\" picture of early-type galaxies. However,  both low-mass Es and high-mass ``dEs\" are now recognized to be quite complex, having been found to contain dust, spiral arms, embedded disks, and bars \\citep{jerjen00, barazza02, derijcke03b, lisker06a, lauraa06}, as well as counter rotating and kinematically decoupled cores \\citep{derijcke04,  thomas06, chilingarian08}, and ongoing star formation \\citep[e.g.,][]{derijcke03a, lisker06b, cote06,michielsen07}. These features suggest that a non-negligible fraction of intermediate-mass galaxies classified as ``early\" types have experienced some level of morphological transformation, likely through mergers, accretions, or interactions with the cluster environment \\citep{moore96, kazantzidis11}.  It is, in fact, possible that the nuclei in some of our early-type galaxies formed in {\\it late-type} progenitors. A recent finding by \\citet{emsellem08} noted that galaxies with S\\'ersic indices of $n\\lesssim3.5$ have compressive tidal forces in their central regions, with the size of the compressive region increasing with decreasing S\\'ersic index. Assuming a constant S\\'ersic index of $n = 1$, the amplitude of the tidal forces was found to scale linearly with galaxy mass, and form a central massive object (CMO) with a constant host mass fraction of $\\sim0.5\\%$. A CMO growing through gas accretion in this way would eventually reach a critical density and luminosity, altering the galaxy profile such that it no longer has central compressive forces. Comparison of this theoretical threshold nucleus luminosity with \\citetalias{cote06} reveals that many observed nuclei are much more luminous than would be predicted by this model, which suggests that the nuclei in early-type galaxies may have formed in some low S\\'ersic index, gas-rich progenitors that have since evolved morphologically.  High-resolution observations of molecular and neutral hydrogen in these galaxies may be able to constrain the role of gas inflow and enrichment in nucleus formation, since H$_{2}$ will highlight regions of star formation, while HI is a tracer of processes affected by the ICM and gravitational interactions. Subarcsecond-resolution  maps of molecular starburst gas --- by using ALMA to observe the CO transitions and EVLA to detect HI through  1.4~GHz emission --- would allow the relationship between galaxy nuclei and molecular gas to be examined in much greater detail than is currently possible.   \\subsubsection{Possible Connections to Black Holes}  As discussed in \\S~\\ref{sec:introduction}, recent observations have uncovered the coexistence of nuclei and black holes in intermediate-mass galaxies, which may have implications for the evolution of the central regions of galaxies. For instance, \\citet{hopkins10b} performed simulations of gas accretion on to a black hole, which they find can form a lopsided, eccentric nuclear disk that exerts a strong torque on and drives in the remainder of the gas, producing a system much like that found in M31.  Another simulation by \\citet{bekki10} examined the merging of two nuclei containing black holes, and found the dynamical heating of the cluster from the black hole binary expelled stars from the center, with the final stellar density of the remnant decreased with increasing black hole mass fraction. This type of merger could produce observed ``core'' galaxies with larger black holes (as originally noted by \\citealt{ebisuzaki91}; see also \\citealt{milosavljevic01}), and shape the inner regions of intermediate-luminosity galaxies in which a nucleus is difficult to distinguish observationally (\\S\\ref{sec:freq}). Their simulations further showed that if only one nucleus had a black hole, the decrease in stellar density of the nucleus was less pronounced, as most of the heating comes from the black hole binary. In mergers where neither nucleus had a black hole, the stellar density of the nucleus increased.  If black holes do become an increasingly dominant component of the CMO mass budget in high- and intermediate-luminosity galaxies, then they could either hinder nucleus growth, or lower the density of the nucleus through mergers until it is destroyed by black hole binary feedback. These effects could create the trends in intermediate-mass galaxy surface brightness profiles observed in this study, where the galaxies undergo a transition from central light ``excesses\" to ``deficits\" as they become more luminous (see also \\citetalias{glass11}).      This HST study examined 43 early-type galaxies in the Fornax cluster, imaged in the ACS F475W and F850LP bands. Our analysis --- performed in both one- and two-dimensions --- extracted photometric and structural parameters for 31 compact stellar nuclei in these early-type galaxies. The main results are summarized as follows:  \\begin{enumerate} \\item We have compared our 1D results to those obtained by using 2D image modeling techniques, and found the extracted nucleus structural parameters to be in agreement for both methods. Although 2D fitting potentially allows for full structural decomposition of a galaxy, 1D methods enable characterization the outer regions with a single surface brightness profile. We conclude that 1D fits are more appropriate for our study, since they allow us to easily compare nucleus and galaxy parameters in an objective and homogeneous way. \\item We find that $72\\pm13$\\% of the 43 galaxies in our sample are nucleated, which is a significant increase from ground-based studies. The nuclei --- defined as a central excess relative to the inward extrapolation of a S\\'ersic model \\citepalias{cote06} --- are found exclusively in galaxies with $M_{B}\\gtrsim-19.5$ (${\\cal M}_* \\lesssim 10^{10.6}{\\cal M}_{\\odot}$),  and the frequency of nucleation for galaxies fainter than this magnitude is $89\\pm16$\\% (31/35). As was found previously in the Virgo cluster, nuclei are exceedingly common in low-mass, early-type galaxies in the Fornax cluster (i.e., ${\\cal M}_* \\gtrsim 10^9{\\cal M}_{\\odot}$). \\item Most nuclei are not significantly offset from their host photocenter --- only three are offset by more than $0\\farcs5$. We do not find any trend between the magnitude of the offset and host galaxy luminosity. \\item We find a nearly constant nucleus-to-galaxy luminosity ratio of $\\approx$0.4\\%. The observed nucleus luminosity function can be understood therefore in terms of the galaxy selection function (and the fact that galaxies brighter than $M_{B}\\lesssim-19.5$ do not contain nuclei). If we parameterize the nucleus luminosity function as a normalized Gaussian, we find peaks at $\\langle M_{g}\\rangle = -11.5$ and $\\langle M_{z}\\rangle = -12.7$~mag, which is $\\sim40$ times more luminous than the peak of the GC luminosity function. The nuclei are also found to have larger sizes and different effective surface brightness scaling relations than the GCs. \\item The colors of the nuclei in hosts with $B_{T}<13.5$ are  found to correlate with galaxy colors, as well as with galaxy and nucleus luminosities. In particular, both the galaxies and the nuclei were observed to become increasingly red with increasing galaxy luminosity, with the trend being steeper for the nuclei. This leads to a relation between nucleus-and-host color difference and host magnitude, where nuclei that are more red than their hosts are found predominantly in brighter galaxies, and vice versa. However, on average most of the nuclei are significantly bluer in $(g-z)$ color than their hosts by $0.28\\pm0.04$ mag. \\item A comparison to \\citetalias{cote06}, which examined the nuclei of early-type galaxies in Virgo, reveals many similarities between the nuclei in the two environments. Both studies find similar frequencies of nucleation (increasing sharply from 0 to $\\gtrsim70$\\% for galaxies with $M_B>-19.5$~mag), surface brightness selection effects, nucleus-to-galaxy luminosity ratios, nucleus luminosity functions, sizes, and color-magnitude relations. The trend along the luminosity function where the galaxy central surface brightness profiles gradually change from having a luminosity ``deficit'' to an ``excess'' is shared by both samples \\citepalias[see also][]{cote07, glass11}, which suggests that generic formation and evolution processes largely independent of the galaxy environment are involved in shaping the central regions of galaxies. Rather, nucleus creation may be more contingent on local factors, especially host galaxy mass. \\end{enumerate}   {\\it Our conclusion is that, in low-mass galaxies, the dominant mechanism for nucleus growth is probably infall of star clusters through dynamical friction, while at higher masses, gas accretion resulting from mergers and torques becomes dominant.} There is no reason to expect either of these processes to be discontinuous, and we argue that the relative importance of these processes vary smoothly as a function of galaxy mass. We examine the efficiency of dynamical friction in our sample galaxies and confirm the finding of many previous studies that  star cluster infall is most effective in low-mass galaxies.  Based on simulations carried out by other researchers, we argue that gas infall, followed by central star formation, becomes increasingly important in high-mass galaxies having S\\'ersic indices that may have been inflated by successive mergers and accretions. There is also some evidence for ``hybrid nuclei\" in some of the intermediate-mass galaxies in our sample: i.e., nuclear components with complex inner structures. Simulations that take into account {\\it multiple} formation mechanisms --- star cluster infall, gas accretion driven by tidal torques and/or accretions and mergers, the influence of central black holes, etc --- are urgently needed to elucidate the processes that drive nucleus formation in different mass regimes.  Both dissipationless cluster infall and gas accretion models make predictions that nucleus formation would depend on local density \\citep{oh00, babul92}. Although the fact that we do not find any major differences between the nuclei of Virgo and Fornax suggests that local density may not be a dominant factor in their formation, observations that examine the entire volume of a galaxy cluster (and that have the sensitivity necessary to detect the nuclei) may help determine the role environment plays in shaping the nuclei and their hosts. In this context, the forthcoming {\\it Next Generation Virgo Cluster Survey} (Ferrarese et~al. 2012), which is imaging the entire Virgo cluster to a (10$\\sigma$) depth of $g \\approx 25.7$, should provide important new constraints on formation models."
    },
    "1208/1208.5940_arXiv.txt": {
        "abstract": "{} {Study of the inner disk structure of dwarf novae (i.e., nonmagnetic cataclysmic variables).} {Power spectral analysis of the X-ray light curves obtained using the Rossi X-ray Timing Explorer ($RXTE$) and  X-ray Multi-mirror Mission (\\XMM) data. We fit such power spectra with a simple model that describes variability due to matter fluctuations. In addition, we perform cross-correlation analysis of simultaneous UV and X-ray light curves using the \\XMM\\ data in order to determine time lags between the different wavelength data.} { We show for five DN systems, SS Cyg, VW Hyi, RU Peg, WW Cet and T Leo that the UV and X-ray power spectra of their time variable light curves are similar in quiescence. All of them show a break in their power spectra, which in the framework of the model of propagating fluctuations indicates inner disk truncation. We derive the inner disk radii for these systems in a range (10-3)$\\times$10$^{9}$ cm. We analyze the $RXTE$ data of SS Cyg in outburst and compare it with the power spectra, obtained during the period of quiescence. We show that during the outburst the disk moves towards the white dwarf and recedes as the outburst declines. We calculate the correlation  between the simultaneous UV and X-ray light curves of the five DN studied in this work, using the \\XMM data obtained in the quiescence and find X-ray time lags of  96-181 sec. This can be explained by the travel time of matter from a truncated inner disk to the white dwarf surface.} { We suggest that, in general, DN may have truncated accretion disks in quiescence which can also explain the UV and X-ray delays in the outburst stage and that the accretion may occur through  coronal flows in the disk (e.g., rotating accretion disk coronae). Within a framework of the model of propagating fluctuations the comparison of the X-ray/UV time lags observed by us in the case of DN systems with those, detected for a magnetic Intermediate Polar allows us to make a rough estimate of the viscosity parameter $\\alpha\\sim0.25$ in the innermost parts of the accretion flow of DN systems.} ",
        "introduction": "Cataclysmic variables (CVs) are interacting binaries hosting a white dwarf (WD) primary star accreting material from a late-type main sequence (MS) star. Accreting material forms a disk that is expected to reach all the way to the WD in cases where the magnetic field of the WD is weak ( $B$ $<$ 0.01 MG) and such systems are referred as nonmagnetic CVs \\cite[see][for a review]{warner95}  In our work we consider dwarf novae -- a subclass of non magnetic cataclysmic variables. In such systems the matter, which is transferred from the secondary star to the Roche lobe of the primary, does not form a stationary flow to the WD surface, but the mass transfer rate is diminishing towards the WD during the so called quiescent state. These states are interrupted every few weeks to tens of years by the enhanced accretion flow (outbursts) that lasts days to weeks and the systems significantly brightens (bolometrically).  The material in the inner disk of dwarf novae initially moving with the Keplerian velocity dissipates its kinetic energy in order to accrete onto the slowly rotating WD creating a boundary layer \\citep{lyndenbell74,kippenhahn78,narayan93,godon95}.  Observations show that during the quiescence (low-mass accretion rates in the innermost parts of the flow) a significant or dominant fraction of the bolometric emission of dwarf novae is emitted via the bremsstrahlung process from an optically thin hot plasma in the hard X-rays \\citep{patterson85,narayan93}. Typical characteristic of the quiescent X-ray emission is a multi-temperature  quasi-isobaric cooling flow model of plasma emission with temperatures of 6-50 keV and an accretion rate of 10$^{-12}$-10$^{-10}$ M$_{\\odot}$/yr with $L_x<10^{33}$ erg s$^{-1}$ \\citep{perna03,pandel05,kuulkers06,rana06,balman11}.  Observational appearance of quiescent dwarf novae indicates that significant part of the energy release during ongoing accretion is happening in the optically thin region near the white dwarf. Thus the accretion flow is likely changing its character from optically thick disk-like flow in the its outer parts to an optically thin \"corona\"-like flow close to the WD.  Observationally, such a truncation of the optically thick accretion disk in dwarf novae in quiescence was invoked due to the observed time lags between the optical and UV fluxes at the rise phase  of the outbursts \\cite[see reviews][]{lasota01,lasota04}, or due to unusual shape of the optical spectra or light curves of DNe \\cite[see e.g.][]{linnell05,kuulkers11}.  Theoretical support for such two-phase flow was given by a model of the disk evaporation of \\cite{meyer94}. This model was later elaborated to show that the disk evaporation (coronal \"syphon\" flow) may create optically thick-optically thin transition regions at various distances from the WD \\citep{liu97,mineshige98}.   Attempts to obtain a map of the accretion disk in cataclysmic variables were done with the help of the eclipse mapping method \\cite[e.g.][]{horne85}, but this method has its own limitations and it is not very sensitive to the innermost parts of the accretion disk.  Recently, an additional diagnostic tool was proposed -  the aperiodic variability of brightness of sources. Virtually all accreting sources demonstrate aperiodic variability of their brightness over a wide range of time scales (see \\citealt{bruch92} for a review of aperiodic variability of cataclysmic variables). While the long time scale variability might be created at the outer parts of the accretion disk \\cite[see e.g.][]{warner71}, the relatively fast time variability (at $f>$few mHz) originates in the inner parts of the accretion flow \\cite[see e.g.][]{bruch00,baptista04}  Properties of this noise is similar to that of the X-ray binaries with neutron stars and black holes. These properties are quite peculiar, which does not allow to explain it as a sum of independent burst-like events in the region of main energy release. In particular, the variability demonstrates very tight relation between the rms amplitude of variations and the average flux level \\cite[see e.g.][]{uttley01,scaringi12}. The variability spans over an extremely wide range of frequencies with the same power-law slopes \\cite[see e.g.][]{churazov01}. Now, the widely accepted model of origin of this aperiodic flicker noise is a model of propagating fluctuations \\citep{lyubarskii97,churazov01,uttley01,arevalo06,revnivtsev09,revnivtsev10,uttley11}.  In the framework of this model, the modulations of the light are created by variations of the instantaneous value of the mass accretion rate in the region of the energy release. These variations of the mass accretion rate, in turn, are inserted into the flow at all radii of the accretion disk due to the stochastic nature of its viscosity and then transferred toward the compact object. Fast variations of the mass accretion rate inserted into the flow at distances closer to the central object modulate the mass flow incoming into these regions from outer parts of the accretion disk.  This model predicts that the truncated accretion disk should lack some part of its variability at high Fourier frequencies, i.e. at the time scales shorter than typical time scale of variability at the inner edge of the disk. This prediction was checked for the accreting systems, in which the disks are indeed truncated due to the interaction with the compact object magnetospheres, in particular - accreting magnetic neutron stars \\citep{revnivtsev09} and accreting magnetic white dwarfs \\citep{revnivtsev10}. The revealed breaks in the power spectra of these accreting binaries allowed one to make estimates of the inner radius of the accretion disk. In particular, in the work of \\cite{revnivtsev11} it was shown that the inner truncation radius of the accretion disk in EX Hya estimated from the variability arguments was quite compatible with those estimated with the help of completely different physical effects.  In this paper we would like to apply the similar diagnostic tool to make estimates of the inner boundary of the optically thick accretion disk in non-magnetic WDs -- dwarf novae. We use \\XMM and $RXTE$ data to study the broad-band noise in DN and  calculate the inner disk radii for five systems, SS Cyg, VW Hyi, RU Peg, WW Cet, and T Leo. Basing on the non-detection of any periodic X-ray light variations (with the very tight upper limit) in flux of the enlisted systems we assume that they contain non-magnetic WDs, while such a classification for SS Cyg is challenged by some authors \\cite[e.g.][]{lombardi87,giovannelli99}. This is not particularly important for our study because we study the accretion disk truncation and the origin of this truncation might have different nature in different sources.  We infer from our results that these systems have truncated disks at large radii and that some form of rotating coronal flow (e.g., ADAF disks and/or accretion disk coronae) exits in DN systems where the material is transported to the surface of the WD. ",
        "conclusions": "We have presented the power spectral analysis of five DNe systems in quiescence and searched for cross-correlations between the X-ray and UV light curves.  We have studied the red noise structure resulting from the flicker noise in the accretion disks and modeled the PDS yielding the break frequencies. This is a strong indication that the optically thick Keplerian flow is truncated at some large radii during the quiescent state of DNe and coronal flows are formed (optically thick-thin disk transition). These structures may be extended on the accretion disk and be emitting at low levels that require  high sensitivity for detection. Such structures may be created by the disk evaporation as suggested earlier. Our range of break frequencies (1-6 mHz) yield a range of radii (10-0.3)$\\times$10$^9$ cm for the inner disk radius with the WD masses 0.4-1.3 M$_{\\odot}$.  In addition, we have also analysed the X-ray outburst data of SS Cyg taken at different epochs and derived that the disk  moves in from the large truncation radius during quiescence towards the WD during the optical peak of the outburst and recedes as the X-rays peak in the outburst, finally to the quiescent inner disk radius.  Our findings are consistent with the previous suggestions of an accretion disk corona existing in SS Cyg.  We modeled the cross-correlations of the quiescent UV and X-ray light curves of our sample of DNe yielding time delays of X-rays in the range 96-181 sec  which indicates the time-lag of emission as the matter travels from the innermost parts of a truncated accretion disk (UV emitting region) to the surface of the WD, where the majority of the X-ray emission is created (i.e., the X-ray photons are delayed). In four of the systems (except for SS Cyg) we also detect zero-time lag correlation indicating the existence of irradiation and reprocessing of X-rays from the cold disk consistent with the light-crossing timescales of the systems.  Finally, in the framework of the propagating fluctuations model we used the ratio of the break timescale and the time-lag of magnetic CV EX Hya and DN systems to derive an estimate of 0.25 for the $\\alpha$ parameter in the inner (optically thin) parts of the accretion flow of DNe disks."
    },
    "1208/1208.5486_arXiv.txt": {
        "abstract": "We present the discovery of three new B[e] supergiants (sgB[e] stars) in the Small Magellanic Cloud (SMC). All three stars (R15, R38, and R48) were identified in the course of our Runaways and Isolated O Type Star Spectroscopic Survey of the SMC (RIOTS4). The stars show optical spectra that closely resemble those of previously known B[e] stars, presenting numerous low-ionization forbidden and permitted emission lines such as [Fe II] and Fe II.  Furthermore, our stars have luminosities of $\\log(L/$L$_{\\odot})$ $\\geq$ 4, demonstrating that they are supergiants.  However, we find lower infrared excesses and weaker forbidden emission lines than for previously identified B[e] supergiants.  Thus our stars appear to either have less material in their circumstellar disks than other sgB[e] stars, or the circumstellar material has lower dust content. We suggest that these may constitute a new subclass of dust-poor sgB[e] stars. ",
        "introduction": "\\label{s:intro}  Massive stars comprise a small fraction of the total number of stars in the universe, and are short-lived. However, they provide a dominant portion of the mechanical and radiative energy of galaxies.  Massive stars can enrich the universe with metals, trigger star formation, and are a potential source for the reionization of the universe.   The evolution of massive stars is central to their ability to influence their surroundings. Therefore, a better understanding of the phases of massive stellar evolution will help us comprehend many physical processes that take place within galaxies.  In massive stars, the B[e] phenomenon is an evolutionary phase whose origin remains unclear.  Stars exhibiting the B[e] phenomenon show numerous emission lines, and normally show infrared radiation from  dust. These properties point to the presence of large amounts of circumstellar material, most likely in a disk configuration.  The B[e] phenomenon is perplexing because stars in many different stages of evolution can exhibit its signs.  \\citet{b:Lamers98} divided the types of stars showing the B[e] phenomenon into five groups: post-main sequence B[e] supergiants (sgB[e] stars), pre-main sequence Herbig B[e] stars, compact planetary nebulae, symbiotic B[e] stars, and an entire category of unclassified B[e] stars.  A link between strong infrared excess and B stars showing forbidden and permitted emission lines was identified by \\citet{b:Geisel70} and \\citet{b:allenswings76}. \\citet{b:Conti76} then suggested that ``B-type stars which show forbidden emission lines'' should be classified as B[e] stars.  Since then, many more B[e] stars have been found, leading to a standard definition of the B[e] phenomenon as follows \\citep[e.g.,][]{b:Zickgraf98, b:Lamers98}: strong Balmer emission lines; low-excitation permitted and forbidden emission lines from low-ionization metals such as Fe; and a strong excess in the near or mid-infrared due to circumstellar dust. Among the more homogeneous and well studied group of stars showing the B[e] phenomenon are the B[e] supergiants.  Supergiant B[e] stars (sgB[e]) also have well defined criteria laid out by \\citet{b:Lamers98}, including supergiant luminosity ($\\log L/L_{\\odot}\\geq$ 4), indicators of evolution off the main sequence sequence, broad ultraviolet (UV) absorption lines, and signs of nitrogen abundance enhancement.  In order to explain the emission properties of sgB[e] stars, \\citet{b:Zickgraf85} proposed a two-component wind model where the star is rotating near break-up velocity and simultaneously close to the Eddington limit.  This causes the star to emit a wind that is fast and sparse near the poles, producing the high excitation absorption, while being denser around the equator, producing the emission features.  The two-component wind allows for the creation of a dusty circumstellar disk, which produces the infrared excess seen around sgB[e] stars.  In this scenario, the star begins exhibiting the B[e] phenomenon sometime after the red supergiant phase, but before He-core burning begins \\citep{b:Lamers06}.  Some sgB[e] stars exhibit properties akin to those of Luminous Blue Variables (LBVs) suggesting that LBVs and at least some sgB[e] stars are related, although their relation remains unclear, as summarized by \\citet[]{b:Lamers06}.  R4 in the Small Magellanic Cloud (SMC) is a prominent candidate for this model.  R4 is currently classified as a LBV/sgB[e] because it shows properties of both objects.  \\citet{b:Pasquali00} further suggested that R4 could have formed in a binary merger.  It is unknown whether this is a common origin for sgB[e] stars in general, or whether R4 is a unique case. Others have suggested that sgB[e] stars arise from interacting binaries, which is consistent with discoveries of binary companions around sgB[e] stars \\citep[e.g.,][]{b:Wheelwright12, b:Miro03}.  Binary companions have been found around other types of B[e] stars as well, such as around Herbig B[e] stars \\citep[e.g.,]{b:Kraus11}.  This scenario could help explain why stars exhibiting the B[e] phenomenon show similar properties while apparently being at different stages of evolution.  In this paper, we report the discovery of three previously unidentified sgB[e] stars in the SMC:  R15, R38, and R48 \\citep[]{b:Feast60}, along with observations of a previously known sgB[e], AzV 154 \\citep[]{b:Azzopardi75}, obtained from the Runaways and Isolated O Star Spectroscopic Survey of the SMC \\citep[RIOTS4;][]{b:Oey11, b:Lamb11}. The RIOTS4 survey is a photometrically selected and spatially complete spectroscopic survey of OB stars in the SMC.  The three previously unknown sgB[e] stars all show optical spectra characteristic of the B[e] phenomenon, including forbidden and permitted emission.  These stars also show additional indicators that they are supergiants, and they appear to be post-main sequence stars. ",
        "conclusions": "\\label{s:conc}  We report spectroscopic observations of three new sgB[e] stars, along with similar observations of AzV 154, a previously known sgB[e] star.  These stars were observed in the course of the RIOTS4 survey of massive field stars in the SMC.  Our stars show rich emission-line spectra, including emission in the Balmer lines, and forbidden and permitted emission from low-excitation metals typical of the B[e] phenomenon.  The stars have supergiant luminosities.  Based on observations of our stars in H$\\alpha$, [SII], and [OIII], they are apparently not associated with star-forming regions, implying that they are unlikely to be Herbig B[e] stars.  The presence of P Cygni profiles in one of the spectra, the comparison to AzV 154, and the position of our stars on the H-R diagram are also consistent with the properties of sgB[e] stars.  Most interestingly, our stars do not show the strong IR excess characteristic of other sgB[e] stars.  Instead, the IR emission from our stars resembles that of classical Be stars, suggesting that it originates from free-free emission.  Thus, our stars lack the strong circumstellar dust disks that produce the IR excess in B[e] stars.  This could be due to low dust content in the disks, or it could be due to the disks being much lower mass. The existence of other dust-poor sgB[e] stars \\citep[e.g.,][]{b:Dunstall12, b:Miroshnichenko98} suggests that such objects are a distinct sub-category of sgB[e] stars, which may be transitioning toward or away from the normal sgB[e] phase."
    },
    "1208/1208.0836.txt": {
        "abstract": "We present 21 examples of C\\,{\\sc iv} Broad Absorption Line (BAL) trough disappearance in 19 quasars selected from systematic multi-epoch observations of 582 bright BAL quasars ($1.9<z<4.5$) by the Sloan Digital Sky Survey-I/II \\hbox{(SDSS-I/II)} and \\hbox{SDSS-III}. The observations span \\hbox{1.1--3.9~yr} rest-frame timescales, longer than have been sampled in many previous BAL variability studies. On these timescales, $\\approx 2.3$\\% of C~{\\sc iv} BAL troughs  disappear and  $\\approx 3.3$\\% of BAL quasars show a disappearing  trough. These observed frequencies suggest that  many C\\,{\\sc iv} BAL absorbers spend on average at most a century along our line of sight to their quasar. Ten of the 19 BAL quasars showing C\\,{\\sc iv} BAL  disappearance have apparently transformed from BAL to non-BAL quasars;  these are the first reported examples of such transformations. The BAL troughs that disappear tend to be those with small-to-moderate equivalent widths, relatively shallow depths, and high outflow velocities. Other non-disappearing C\\,{\\sc iv} BALs in those nine objects having multiple troughs tend  to weaken when one of them disappears, indicating a connection between the disappearing and non-disappearing troughs, even for velocity separations as large as \\hbox{10000--15000 $\\mathrm{km\\,s^{-1}}$}. We discuss possible origins of this connection including disk-wind rotation and changes in shielding gas. ",
        "introduction": "\\label{intro}  Intrinsic absorption lines in quasar spectra are often produced by outflowing winds along the line of sight that are launched from the accretion disk or other structures around the central supermassive black hole \\citep[SMBH; e.g.,][]{murray95,proga00}. Such absorption lines, shaped by the geometry and kinematic structure of the outflows, appear in quasar spectra as  broad absorption lines (BALs; $\\Delta v$\\,$\\geq$\\,2000\\,$\\mathrm{km\\,s^{-1}}$), mini-BALs (2000\\,$\\geq$\\,$\\Delta v$\\,$\\geq$\\,500\\,$\\mathrm{km\\,s^{-1}}$), or intrinsic narrow absorption lines (NALs; $\\Delta v$\\,$\\leq$\\,500\\,$\\mathrm{km\\,s^{-1}}$); e.g., see \\citet{hamann08} and \\citet{gb08}. Traditionally defined BAL troughs are observed in $\\approx 15$\\% of quasars; these troughs are sufficiently strong that the depths of the features lie at least 10\\% under the continuum in the, e.g., Si\\,{\\sc iv}~$\\lambda$\\,1400, C\\,{\\sc iv}~$\\lambda$\\,1549, Al\\,{\\sc iii}~$\\lambda$\\,1857, or Mg\\,{\\sc ii}~$\\lambda$\\,2799 transitions with blueshifted velocities up to $\\approx 0.1c$ (e.g., \\citealp{wey91}; \\citealp{gibson09}, hereafter G09; \\citealp{allen11}; and references therein). Quasars that present BAL troughs in their spectra are often classified into three subtypes depending on the presence of absorption lines in specified transitions: (1) High-ionization BAL quasars show absorption lines in high-ionization transitions including Si\\,{\\sc iv}, and C\\,{\\sc iv}. (2) Low-ionization BAL quasars possess Mg\\,{\\sc ii} and/or Al\\,{\\sc iii} absorption lines, in addition to the high-ionization transitions. (3) Iron low-ionization BAL quasars show additional absorption from excited states of Fe\\,{\\sc ii} and Fe\\,{\\sc iii} \\citep[e.g.,][and references therein]{hall02}.  The winds revealed by intrinsic quasar absorption lines are of importance for two main reasons. First, the observed frequency of quasar absorption lines indicates that these winds are a significant component of the nuclear environment. Indeed, disk accretion onto the SMBH may require significant mass ejection for expulsion of angular momentum from the system \\citep[e.g.,][]{blandford82,crenshaw03}. Second, quasar winds may be agents of feedback into massive galaxies, regulating star formation and further SMBH accretion via the removal of cold gas \\citep[e.g.,][]{springel05,king10}.  BAL troughs, the strongest observed signatures of quasar winds, often vary in equivalent width (EW) and/or shape over rest-frame timescales of months to years \\citep[e.g.,][]{barlow93,lundgren07,gibson08,gibson10,cap11,cap12}. Recent statistical studies of BAL variability have shown that the fractional EW change increases with rest-frame timescale over the range \\hbox{0.05--5~yr}. Such variations could, in principle, be driven by changes in covering factor, velocity structure, or ionization level. Of these possibilities, the generally favored driver for most BAL variations is changes in the covering factor of outflow stream lines that partially block the continuum emission \\citep[e.g.,][]{hamann98, arav99}. These covering-factor changes could ultimately be caused by rotation of a non-axisymmetric outflow that is loosely anchored to the accretion disk, or they could arise from large-scale changes in wind structure \\citep[e.g.,][]{proga00}. Changes in ionization level are generally disfavored as a primary driver, since BAL troughs are often highly saturated and thus should be only weakly responsive to ionization-level changes. Furthermore, BAL trough variations generally do not appear to be correlated with variations of the observable continuum (typically) longward of Ly$\\alpha$ \\citep[e.g.,][]{gibson08}, although the observable continuum is not the ionizing continuum for the BAL gas (i.e., it is possible that the two may vary differently).   Strong variations in BAL EWs observed over multi-year timescales in the rest frame suggest that BAL disappearance and BAL emergence may be significant effects on such timescales \\citep[e.g.,][]{gibson08,hall11}. However, largely owing to practical difficulties in observing large samples of BAL quasars on multi-year timescales (often corresponding to \\hbox{10--20~yr} in the observed frame),  only a small number of BAL disappearance \\citep{junk01,lundgren07,hall11,vivek12} and emergence \\citep{ma02,lundgren07,hamann08,leighly09,krongold10,cap12,vivek12} events have been discovered. Specifically, considering the BAL disappearance phenomenon of most relevance to this paper, we are only aware of four reported cases. In the first, a Mg\\,{\\sc ii} BAL in the spectrum of the binary quasar \\hbox{LBQS~0103--2753}  essentially disappeared over \\hbox{$\\leq 6.0$} rest-frame years, converting this object from a low-ionization BAL quasar to a high-ionization BAL quasar \\citep{junk01}.% In the second, the highest velocity C\\,{\\sc iv} trough in the spectrum of J075010.17+304032.3 nearly disappeared in \\hbox{$\\leq$ 0.3} rest-frame years, which is the only known example of C\\,{\\sc iv} disappearance \\citep{lundgren07}. In the third, the  Fe\\,{\\sc ii} troughs in the spectrum of FBQS~J1408+3054 disappeared over \\hbox{$\\leq$ 5.1} rest-frame years, converting this object from an iron low-ionization BAL quasar to a low-ionization BAL quasar \\citep{hall11}. In the fourth, a Mg\\,{\\sc ii} BAL trough in the spectrum of SDSS J133356.02+001229.1 disappeared in \\hbox{$\\leq$ 3.7} rest-frame years and another Mg\\,{\\sc ii} BAL trough --- that emerged at higher velocity --- nearly disappeared in \\hbox{$\\leq$ 5.7} rest-frame years \\citep{vivek12}. These disappearance examples did not involve transformations from BAL quasars to non-BAL quasar status because troughs from other ions, or additional C\\,{\\sc iv} troughs, remained.   We have compiled a sample of 19 quasars with disappearing BAL troughs from a well-defined, large sample of BAL quasars whose spectra were observed over \\hbox{1.1--3.9~yr} in the rest frame. We aim to define the basic statistical characteristics of the BAL disappearance phenomenon. % In this study, we will discuss the disappearance of BAL troughs as well as the transformation of BAL quasars to non-BAL quasars. These two phenomena can be distinct given that some BAL quasars have multiple troughs. A quasar that possesses more than one BAL trough could have only one of them disappear without transforming to a non-BAL quasar. We only consider BAL quasars without any remaining BAL troughs at a later epoch (in any transition) to have transformed from a BAL quasar to a non-BAL quasar. % The observations and the selection of the main sample used in this study are discussed in \\S2. Identification of disappearing BAL troughs is explained in \\S3. Statistical results and comparisons are presented in \\S4. A summary of the main results and some future prospects are given in \\S5.  Throughout this work we use a cosmology in which $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M=0.3$, and $\\Omega_{\\Lambda}=0.7$. All timescales are in the rest frame unless otherwise mentioned. ",
        "conclusions": "\\label{summary}  We have used a systematically observed sample of 582 BAL quasars with 925 distinct C\\,{\\sc iv} BAL troughs to provide the first statistically meaningful constraints upon BAL disappearance on multi-year timescales. Our main results are the following:  \\begin{enumerate}  \\item We have identified 21 cases of C\\,{\\sc iv} BAL disappearance in 19 quasars. On rest-frame timescales of \\hbox{1.1--3.9~yr}, $f_{\\rm disappear} = 2.3^{+0.6}_{-0.5}$\\% of BAL troughs disappear and $f_{\\rm quasar} = 3.3^{+0.9}_{-0.7}$\\% of BAL quasars show a disappearing trough. If we consider only the pristine sample defined in \\S \\ref{selection}, then we find 11 cases of C\\,{\\sc iv} BAL disappearance in 11 quasars; the corresponding percentages are $f_{\\rm disappear} = 1.2^{+0.4}_{-0.4}$\\% and $f_{\\rm quasar} = 1.9^{+0.8}_{-0.6}$\\%. See \\S \\ref{common}.  \\item The observed frequency of disappearing C\\,{\\sc iv} BAL troughs suggests an average trough rest-frame lifetime of 100--200 yr. See \\S \\ref{common}.  \\item Ten quasars showing C\\,{\\sc iv} BAL disappearance have apparently transformed from BAL to non-BAL quasars;  these are the first reported examples of such transformations. The frequency of BAL to non-BAL quasar transformation on timescales of \\hbox{1.1--3.9~yr} is $1.7^{+0.7}_{-0.5} \\%$. See \\S \\ref{common}.  \\item The BAL quasars with disappearing troughs have representative luminosities ($M_i$ values), SMBH mass estimates, and intrinsic reddening compared to our sample as a whole. See \\S \\ref{qsos}.  \\item As expected from the fact that most BAL quasars are radio quiet, most BAL quasars with disappearing troughs are radio quiet. However, we do find one such quasar that is radio loud and another that is radio intermediate. Thus, BAL disappearance is a phenomenon of both radio-quiet  and radio-loud quasars. See \\S \\ref{qsos}.  \\item BAL disappearance appears to occur mainly for weak or moderate-strength absorption troughs but not the strongest ones; e.g., no troughs with EW~$>12$~\\AA\\ disappeared. The BAL troughs that disappear are shallower than BAL troughs in general, although some fairly deep BAL troughs do disappear. See \\S \\ref{ews}.  \\item Disappearing C\\,{\\sc iv} BAL troughs show higher outflow velocities than BAL troughs in general, as indicated by their measured central velocities and minimum velocities (though their measured maximum velocities do not appear exceptional). There is also suggestive evidence that disappearing BAL troughs tend to be narrower than BAL troughs in general. This tendency for BAL disappearance to occur for higher velocity troughs could be related to, or even a secondary effect of, the fact that BAL disappearance appears to occur mainly for weak or moderate-strength absorption troughs (see point~5 above). See  \\S \\ref{ews}.   \\item When one BAL trough in a quasar spectrum disappears, the other present troughs usually weaken (11 times out of 12 in our sample, corresponding to a significance level $>~99$\\%). The phenomenon occurs even for velocity offsets as large as \\hbox{10000--15000 $\\mathrm{km\\,s^{-1}}$}. Variability across multiple troughs appears surprisingly coordinated. Possible causes of such coordinated variations could be disk-wind rotation or variations of shielding gas that lead to variations of ionizing-continuum radiation. These possible agents will need to be considered in future models of quasar winds. See \\S \\ref{connection}.  \\end{enumerate}  Given the results above, we can identify several promising observational projects that should extend understanding of BAL disappearance. Further spectroscopy of the quasars that have shown BAL disappearance will allow a search for reappearance of any of these BALs. Such reappearances at the same measured velocities would not be expected if wind stream lines have moved out of the line of sight owing to rotation of a non-axisymmetric outflow. However, BAL reappearance should be possible if the disappearance is a consequence of BAL weakening to strengths below our detection threshold. Systematic large-sample variability studies should let us assess the extent to which BAL disappearance is just the extension of normal BAL variability down to very small EWs. Further spectroscopy will also allow monitoring of the additional non-disappearing troughs. Furthermore, the planned absolute flux calibration of the BOSS spectra  \\citep[e.g.,][]{margala11} will allow a search for any systematic continuum-level changes associated with BAL disappearance. The rate of BAL emergence events must balance that of BAL disappearance events if the BAL quasar population is in a steady state, and thus systematic large-scale studies of BAL emergence will be a critical complement to those of disappearance. Finally, multiwavelength observations of the quasars showing BAL disappearance are worthwhile. For example, \\hbox{X-ray} observations of objects that have transformed from BAL to non-BAL quasars will be able to assess if the X-ray absorbing shielding gas is still present along the line of sight.  The main-sample data set utilized in this study, along with the still incoming BOSS observations, will be effective for a variety of additional investigations of BAL variability. These include studies of (1) absorption EW variability as a function of timescale for different BAL transitions, (2) connections between BAL, emission-line, and reddening variability, and (3) the effects of luminosity, redshift, SMBH mass, Eddington fraction, and radio properties on BAL variability. %We intend to investigate these issues in future work, hopefully leading to %insights about the launching of and feedback from quasar winds."
    },
    "1208/1208.3119_arXiv.txt": {
        "abstract": "We use a dense, complete redshift survey, the Smithsonian Hectospec Lensing Survey (SHELS), covering a 4 square degree region of a deep imaging survey, the Deep Lens Survey (DLS), to study the optical spectral properties of Wide-field Infrared Survey Explorer ({\\it WISE}) 22 $\\mu$m-selected galaxies. Among 507 \\wise 22 $\\mu$m-selected sources with (S/N)$_{22\\mu{\\rm m}}\\geq3$ ($\\approx S_{22\\mu m}\\gtrsim2.5$ mJy), we identify the optical counterparts of 481 sources ($\\sim98\\%$) at $R<25.2$ in the very deep, DLS $R$-band source catalog. Among them, 337 galaxies at $R<21$ have SHELS spectroscopic data. Most of these objects are at $z<0.8$. The infrared (IR) luminosities are in the range $4.5\\times10^8  ({\\rm L}_\\odot) \\lesssim L_{IR} \\lesssim 5.4\\times10^{12}  ({\\rm L}_\\odot)$. Most 22 $\\mu$m-selected galaxies are dusty star-forming galaxies with a small ($<$1.5) 4000 \\AA~break. The stacked spectra of the 22 $\\mu$m-selected galaxies binned in IR luminosity show that the strength of the [O III] line relative to H$\\beta$ grows with increasing IR luminosity. The optical spectra of the 22 $\\mu$m-selected galaxies also show that there are some ($\\sim2.8\\%$) unusual galaxies with very strong [Ne III] $\\lambda$3869, 3968 emission lines that require hard ionizing radiation such as AGN or extremely young massive stars. The specific star formation rates (sSFRs) derived from the 3.6 and 22 $\\mu$m flux densities are enhanced if the 22 $\\mu$m-selected galaxies have close late-type neighbors. The sSFR distribution of the 22 $\\mu$m-selected galaxies containing active galactic nuclei (AGNs) is similar to the distribution for star-forming galaxies without AGNs. We identify 48 dust-obscured galaxy (DOG) candidates with large ($\\gtrsim1000$) mid-IR to optical flux density ratio. The combination of deep photometric and spectroscopic data with \\wise data suggests that \\wise can probe the universe to $z\\sim2$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.4163_arXiv.txt": {
        "abstract": "We investigate circumstances under which one can generalize Horndeski's most general scalar-tensor theory of gravity.  Specifically we demonstrate that a nonlinear combination of purely kinetic gravity terms can give rise to an accelerating universe without the addition of extra propagating degrees of freedom on cosmological backgrounds, and exhibit self tuning to bring a large cosmological constant under control.  This nonlinear approach leads to new properties that may be instructive for exploring the behaviors of gravity. ",
        "introduction": "Theories of gravitation on cosmic scales have become an area of intense interest, both as a possible explanation for the observed cosmic acceleration and as an exploration of consistent extensions of general relativity.  Generically such extensions lead to additional degrees of freedom, e.g.\\ scalar modes in scalar-tensor theories, with possible pitfalls of higher than second order derivative field equations that may lack a well posed initial value formulation, or of ghosts and other instabilities.  Horndeski in 1974 wrote the most general scalar-tensor theory giving second order field equations in four dimensional spacetimes \\cite{horndeski}. In an alternate view, Galileon theories \\cite{nicolis,deffayet1,deffayet2}, shift symmetric scalar fields possessing nonlinear combinations of field derivatives, have recently been studied with interest as sound models capable of cosmic acceleration. Also recently, as an approach to solve the cosmological constant problem, a linear combination of four terms called the Fab Four has been identified \\cite{fab4,fab4long,fab4new} as the unique terms allowing self tuning vacua that can cancel a large bare $\\Lambda$ term.  Here we draw on aspects of all three of these approaches to demonstrate that nonlinear combinations of terms involving shift symmetric scalar fields can possess interesting advantages and properties.  This extension of the ``most general'' scalar-tensor theory retains second order field equations and avoids pathologies on symmetric spacetimes such as the usual Friedmann-Lema\\^\\i tre-Robertson-Walker (FLRW) and de Sitter cosmologies. While the nonlinear approach can be applied quite generally, we give a proof of principle using a simple example of purely kinetic couplings with noncanonical forms, extending the ``purely kinetic gravity'' of \\cite{gub}.  Our example employs a nonlinear combination of the standard kinetic term and derivative coupling to the Einstein tensor. Besides possessing second order field equations it does not add any further propagating degrees of freedom, and it can achieve lasting cosmic acceleration unlike the linear, canonical, purely kinetic gravity theory of \\cite{gub}, avoid at least some instabilities unlike the derivatively coupled Galileon theory investigated by \\cite{applin}, and self tune away a cosmological constant like the Fab Four.  Deeper implications exist beyond our simple proof of principle. We emphasize that the example given is intended purely as a proof of principle to inspire further investigation into the theoretical properties of general nonlinear combinations, and not as a fit to observations.  In Sec.~\\ref{sec:method} we explain our nonlinear generalization procedure and the conditions under which no additional propagating degrees of freedom are generated.  The equations of motion are solved in Sec.~\\ref{sec:eom} on a FLRW background, giving the cosmic and field evolution complete with attractors, revealing two distinct ways of approaching a de Sitter asymptotic state. Section~\\ref{sec:tune} demonstrates the self tuning properties of the theory, erasing an initial cosmological constant.  The perturbed equations in Sec.~\\ref{sec:perturb} yield the no-ghost and stability conditions and the evolution of the effective Newton's constant $\\geff$.  We discuss various implications of the results in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  Gravitation is a fundamental force that we have just begun to explore cosmologically.  One of the great advances made in gravity research in the past few years is the realization that symmetry principles both strongly restrict the theory and open up new avenues and effects.  Galileon gravity and massive gravity both use shift symmetric fields and their couplings to functions of the metric to enable new properties, including cosmic acceleration without a cosmological constant or field potential. An action allowed by the symmetries and well behaved in initial value formulation, specifically one leading to second order equations of motion, is of particular interest.  If moreover the field exhibits self tuning, allowing it to overcome a high energy cosmological constant, the theory is well worth examining.  We show that by promoting a purely kinetic gravity term to a nonlinear function, possibly mixed with a noncanonical kinetic term for the field, fascinating properties can ensue.  In addition to second order equations of motion and self tuning, the theory does not incur extra propagating degrees of freedom on a cosmological background. Similar effects of symmetric backgrounds are seen in massive gravity. For example, in massive gravity it was found \\cite{massiveG} that in isotropic spacetimes the shift symmetric (St\\\"uckelberg) fields have vanishing mixing between the graviton and the scalar mode, and furthermore their kinetic terms vanish.  The new term discussed here, ``Fab 5 Freddy,'' can self accelerate and is merely the harbinger of a whole class of such nonlinear promotions or combinations of terms.  The background evolution of the expansion and field lead to early time tracker behavior and late time de Sitter attractors.  Solving the linear perturbation equations we see that simple power law functions can be free of ghosts.  The gravitational coupling and dispersion relation of perturbations become scale dependent, possibly leading to an early time instability and a late time vanishing of gravity.  The specific models studied may not be observationally viable but the characteristics arising from the nonlinear, noncanonical action open new aspects of gravity. Most intriguing is the self tuning property that can cancel a bare cosmological constant dynamically, even through phase transitions. The evolution of the field basically makes $\\Lambda$ invisible.  That the most general scalar-tensor theory giving second order field equations in 4D could be further generalized, at least on cosmological backgrounds, is highly interesting. The specific term considered here, a nonlinear promotion of the field kinetics coupled to the Einstein tensor (the unique, low mass dimension shift symmetric combination giving second order field equations), is merely a proof of principle, while theoretically instructive.  Ways to extend this class of theory more generally, to different nonlinear functions and combinations, are straightforward and may preserve the most interesting and desirable characteristics while leading to more viable predictions experimentally."
    },
    "1208/1208.0356_arXiv.txt": {
        "abstract": "The equilibrium theory of the 2D magnetohydrodynamic equations is derived, accounting for the full infinite hierarchies of conserved integrals. An exact description in terms of two coupled elastic membranes emerges, producing long-ranged correlations between the magnetic and velocity fields. This is quite different from the results of previous variational treatments, which relied on a local product ansatz for the thermodynamic Gibbs distribution. The equilibria display the same type of coherent structures, such as compact eddies and zonal jets, previously found in pure fluid equilibria. Possible consequences of this for recent simulations of the solar tachocline are discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0160_arXiv.txt": {
        "abstract": "{ The transient, ultra-luminous X-ray source CXOM31~J004253.1+411422 in the Andromeda galaxy is most likely a 10 solar mass black hole, with super-Eddington luminosity at its peak. The {\\it XMM-Newton} spectra taken during the decline then trace luminosities of $0.86-0.27 \\LEdd$. These spectra are all dominated by a hot disc component, which roughly follows a constant inner radius track in luminosity and temperature as the source declines. At the highest luminosity the disc structure should change due to advection of radiation through the disc. This advected flux can be partly released at lower radii thus modifying the spectral shape. To study the effect of advection at luminosities close to Eddington we employ a fully relativistic slim disc model, {\\tt SLIMBH}, that includes advective cooling and full radiative transfer through the photosphere based on {\\sc TLUSTY}. The model also incorporates relativistic photon ray-tracing from the proper location of the disc photosphere rather than the mid-plane as the slim disc is no longer geometrically thin. We find that these new models differ only slightly from the non-advective (standard) {\\tt BHSPEC} models even at the highest luminosities considered here. While both discs can fit the highest luminosity data, neither is a very good fit to the {\\it lower} luminosities. This could indicate a missing physical process that acts in low luminosity discs and subsides as the disc luminosity approaches the Eddington limit. } \\authorrunning{O.\\,Straub, C. Done \\& M. Middleton} \\titlerunning{Advection in spectra of M31 ULX-1} ",
        "introduction": "\\label{sec:intro} The most luminous objects in the Universe are powered by accretion of matter onto compact objects like neutron stars and black holes. The energy released depends on the mass and spin of the central black hole as well as the mass accretion rate through the disc, but there is a maximum luminosity for which gravity is able to balance the outward pressure of radiation. Ultra-luminous X-ray sources (ULX) are defined as objects whose bolometric luminosity exceeds this Eddington limit ($L > 10^{39} erg \\,\\, s^{-1}$). These objects could be powered by normal (sub-Eddington) accretion onto intermediate mass black holes ($M_{BH} \\simeq 10^2-10^5 \\Msun$). However, there are theoretical problems concerning the formation of such objects \\citep{kin+01}, so these are unlikely to form the bulk of the ULX population. Instead, the majority of sources are probably {\\it super}-Eddington accreting stellar mass black holes. Such flows come in two flavours, either exceeding the Eddington limit by powering strong outflows \\citep{sha+73, lip99, pou+07}, or by advecting the radiation along with the flow (Polish doughnuts: \\citet{abr+78,jar+80}; slim discs: \\citet{abr+88, sad09}). These processes can occur together \\citep{pou+07}, as shown in recent numerical simulations of super-Eddington flows \\citep{ohs+09, ohs+11}. In all these cases, radiation is emitted at the local Eddington limit, so that the total source luminosity is $L \\sim \\LEdd (1 + \\ln \\Mdot/\\MdotEdd)$. Advective accretion discs that can locally exceed the Eddington limit and thus excite winds were proposed by \\citet{dot+11}.  The {\\tt BHSPEC} models are the best current sub-Eddington accretion flow models for stellar mass black holes. Those calculate the spectrum from full radiative transfer through a disc atmosphere for a Novikov-Thorne (general relativistic) emissivity, including ray-tracing that comprises all general and special relativistic effects \\citep{dav+05}. With increasing mass accretion rate, however, the disc's ability to radiate the energy dissipated by viscous stresses ceases. Consequently, the disc overheats and inflates, and the Novikov-Thorne model breaks down. The slim disc is a stable solution for trans-Eddington accretion flows. It is cooled by advection that sweeps some of the emitted energy along with the flow as photons are trapped in the optically thick disc. The photons can be released again at lower radii as the material accelerates towards the black hole. The {\\tt SLIMBH} models, a modification of {\\tt BHSPEC}, incorporate the effects of advection, such as the increasing height of the photosphere and the shift of the inner disc edge towards smaller radii \\citep{sad+11}. {\\tt SLIMBH} models were recently fitted to spectra from the accreting black hole in LMC X-3. However, the differences between standard thin disc models and {\\tt SLIMBH} are not large there although this object reaches $\\sim 0.6 \\LEdd$ \\citep{str+11}.  Here instead we fit these new models to the higher Eddington fraction flows seen in ULX. In particular, we use the XMM-Newton spectrum from the transient ULX CXOM31 J004253.1+411422 (hereafter M31 ULX-1). The advantage of this object is that it showed a steady exponential decline in luminosity in five XMM-Newton spectra from $1.6 - 0.8 \\times 10^{39} erg \\,\\, s^{-1}$ after its discovery at $5 \\times 10^{39} erg \\,\\, s^{-1}$ in Chandra imaging data. Such exponential decays with this time scale are well known from transient low mass X-ray binary black holes in our Galaxy, making it most probable that this ULX is also a $\\sim10 \\Msun$ black hole \\citep{mid+12, kau+12}. This mass would then imply that the luminosity sampled by the XMM-Newton spectra corresponds to $0.86-0.27 \\LEdd$, considerably higher than seen in LMC X-3. Thus it provides an ideal testing ground for the slim disc models, especially as the Galactic absorption column is fairly low, allowing the broad band disc shape to be seen down to low energies.  The paper is structured as follows. The origin and analysis of our data is specified in section~\\ref{sec:data}, the modelling of the ULX spectra is described in section~\\ref{sec:modelling} and the results are discussed in section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} It has long been assumed that in the trans-Eddington luminosity regime two effects could become increasingly important for accretion disc models, (i) advection \\citep{abr+88,min+00} and/or (ii) outflows \\citep{sha+73, pou+07, ohs+09}. In this paper we pin down the effect of advection by fitting ULX spectra with slim discs. We find that advection typically changes the best fitting spectral shape by removing photons in the intermediate energy range where the spectral peak is located and by releasing more photons at high energies (see in Figure~\\ref{fig:comp_ratio}, upper panel). This is the characteristic signature of slim discs where photons are trapped in the accretion flow, accreted inward and released closer to the black hole and even inside the plunging region (see \\citet{abr+10} for a discussion of the inner edge of slim discs and \\citet{zhu+12} for the relevance of radiation emitted from inside the plunging region).  However, these changes make very little difference to the emitted spectrum for $L < \\LEdd$, and we show here that this cannot be the reason for the mismatch between data and models seen in moderate luminosity ($0.1-0.5 \\LEdd$) spectra from both this ULX and from galactic black hole binaries \\citep{kol+11,str+13}. Real accretion discs at these moderate luminosities have a spectrum which is subtly broader than is currently expected from the best disc models. This additional broadening is most marked at lower mass accretion rates and gets weaker as the disc luminosity increases. This is exactly opposite to the luminosity behaviour expected from either winds, advection, bulk turbulence and/or inhomogeneities in the disc.  Instead, it could be connected to the (low level) coronal emission. We cannot constrain a weak high energy tail in our data, but this could potentially illuminate the disc and hence change the structure of the photosphere. However, we note that the most disc dominated GX339-4 spectrum in \\citet{kol+11}, where the tail carries less than a few per cent of the total bolometric flux, shows this additional broadening. This is a very low level of illumination to cause a noticeable change in disc structure. Self-illumination of the inner disc by light-bending around the black hole may be a larger effect when the coronal emission is low \\citep[e.g.,][]{min+04}. However, this should involve a constant fraction of the emission, rather than give a decreasing fraction as $L/\\LEdd$ increases.  Here we speculate that the broadened spectral shape at low luminosity may actually be a signature of magnetic pressure support in the disc. This will increase the scale height of the disc over and above that expected from a standard disc, so decreasing its density and hence its true absorption opacity. Electron scattering then becomes more important, broadening the spectrum \\citep{dav+09}. The extent of this change in disc structure depends on the ratio of magnetic pressure to total (radiation plus gas plus magnetic) pressure. This could explain the observed trend in the data if the magnetic pressure saturates to some fraction of the gas pressure rather than total pressure, as it would decrease in importance as $L/\\LEdd$ increases.  Other observations also strongly indicate that there is a change in behaviour in magnetic stress scaling between the gas pressure and radiation pressure dominated regimes. The rapid rise to outburst, which occurs (mostly) in the gas pressure dominated regime, requires very efficient transport of angular momentum as parametrised by a Shakura-Sunyaev $\\alpha$ viscosity of $\\sim 0.1$ \\citep{dub+99}. However, if this efficient transport is maintained into the radiation pressure dominated regime then the disc becomes effectively optically thin, and its colour temperature correction increases markedly. This is in sharp contrast to the data, where the observed constancy of the colour temperature correction requires $\\alpha<0.01$ at high $L/\\LEdd$, where the disc is in the radiation pressure dominated regime \\citep{don+08, str+11}. Thus the data from the ULX reported here, and previous work on the disc spectra seen from black hole binaries both support a change in magnetic stress scaling with total pressure. Current numerical simulations do not show this behaviour \\citep{hir+09}, but these are at the limit of computational capabilities. Future research on how the self generated magnetic field from the MRI scales with pressure could resolve this issue."
    },
    "1208/1208.2954_arXiv.txt": {
        "abstract": "We present a method for performing data-driven simulations of solar active region formation and evolution. The approach is based on magnetofriction, which evolves the induction equation assuming the plasma velocity is proportional to the Lorentz force. The simulations of active region coronal field are driven by temporal sequences of photospheric magnetograms from the Helioseismic Magnetic Imager (HMI) instrument onboard the Solar Dynamics Observatory (SDO). Under certain conditions, the data-driven simulations produce flux ropes that are ejected from the modeled active region due to loss of equilibrium. Following the ejection of flux ropes, we find an enhancement of the photospheric horizontal field near the polarity inversion line. We also present a method for the synthesis of mock coronal images based on a proxy emissivity calculated from the current density distribution in the model. This method yields mock coronal images that are somewhat reminiscent of images of active regions taken by instruments such as SDO's Atmospheric Imaging Assembly (AIA) at extreme ultraviolet wavelengths. ",
        "introduction": "Understanding the structure and evolution of the solar coronal magnetic field is an important component in understanding space weather.  These dynamics are a consequence of the transport of energy through the photosphere and into the chromosphere and corona, and of the subsequent reorganization by the coronal magnetic field in response to this energy input.  Some of this energy is used to accelerate the solar wind, some is used to heat plasma trapped in closed coronal loops (after which it is emitted as radiation), and some is used to power flares and eruptive phenomena as coronal mass ejections (CMEs).  Coronae above active regions have received particular attention~\\citep[e.g.,][]{RegnierPriest:2007, Kazachenko:2012} because these are the sites that produce the most numerous and strongest flares, and are most often associated with CMEs and particle acceleration events.  Additionally, active region coronae are copious emitters in the extreme ultraviolet (EUV) and x-ray wavelengths of light, with emissions typically taking the form of loop-shaped structures that are assumed to trace out the trajectories of coronal magnetic field lines.  Despite considerable progress in observing the structure and evolution of the solar corona, the root causes of many phenomena remain elusive.  Much of this uncertainty stems from the fact that it is difficult to accurately map out the three-dimensional geometry of the coronal magnetic field, and to observe its temporal evolution.  Direct measurements of the coronal field are sometimes available off the solar limb~\\citep[e.g.,][]{Lin:CoronalFieldMeasurements,BrosiusWhite:2006, Tomczyk:CoronalPolarization,Liu:CoronalIR}, but these measurements often lack sufficient spatial or temporal resolution, and suffer from line-of-sight confusion, thereby making proper interpretation somewhat challenging.  As a result of these difficulties, much effort has been put toward modeling the coronal magnetic field using photospheric magnetograms and/or coronal loop imagery as constraints.  Many models make use of photospheric measurements of the magnetic field as boundary conditions.  These models include potential (current-free) field extrapolations which provide a general idea of the global connectivity, force-free models which accommodate currents but assume the corona is in static force balance, and more physically realistic magnetohydrodynamic (MHD) models which capture many more of the important dynamical processes.  Of these various coronal magnetic field models, only the potential-field source-surface (PFSS) model can presently be computed in real time in a forecasting (predictive) capacity. However, because the active region coronae of particular interest are those that likely contain significant currents, the applicability of the PFSS model for understanding energy buildup and release in active regions is limited, restricting the PFSS model to instances where only visualizing the large-scale (global) geometry of the coronal field, or providing contextual magnetic fields surrounding active regions during relatively quiescent periods~\\citep{Riley:2006}, is needed.  At the other end of the spectrum of coronal field models are MHD simulations ~\\citep[e.g.,][]{Mikic:1999}. These simulations are the most sophisticated in terms of physical realism, but, even with today's supercomputing technology, remain computationally expensive for active-region-sized domains at appreciable resolution.  Additionally, because not all of the necessary photospheric boundary conditions (e.g., gas pressures and electric fields) are measured directly, and because they still require parameterized models to bypass treatment of the microphysics (e.g., coronal heating functions) which are not always well known, even MHD models are not free from assumptions.  In between the PFSS and MHD models lie the class of nonlinear force-free field (NLFFF) models of the corona. NLFFF models contain no dynamics, but do allow for field-aligned electric currents and thus enable measurements of magnetic free energies and magnetic helicity. Fast algorithms that extrapolate the magnetic field upward into the corona have been developed, though with mixed results in terms of their ability to reliably reproduce the observed coronal field structures~\\citep{Schrijver:NLFFF,DeRosa:NLFFF}.  Yet, because the electric current systems that occupy the corona tend to be isolated in and around active regions, NLFFF models seem particularly well suited for detailed studies of the build-up and release of energy within the corona.  One NLFFF solution technique, the magnetofrictional (hereafter MF) scheme~\\citep{YangSturrockAntiochos:Magnetofriction,CraigSneyd:Magnetofriction}, involves integrating a model forward in such a way as to reduce the magnetic stress in the model.  This is achieved by assuming that the inductive velocity is parallel to the Lorentz force, which over time causes the magnetic field in the model to relax to a force-free state.  This method has previously been employed as a way to create a NLFFF from a non-force-free initialization, either via the specification of the vector field on the lower boundary of the domain~\\citep[][]{Valori:NLFFFExtrapolation,Valori:NLFFFExtrapolationII} or via the insertion of a flux-rope structure into an otherwise potential field~\\citep[e.g.,][]{vanBallegooijen:2004,Bobra:FluxRopeInsertion, Savcheva:2009,Savcheva:2012}.   Alternatively, one may drive a MF model by allowing the lower-boundary data to change with time.  Long-term studies of the formation and subsequent evolution of filament channels using MF schemes have been performed by~\\citet*{vanBallegooijenPriestMackay:MeanFieldModel} and~\\citet*{Mackay:FilamentChirality}, in which it was found that potential-field models did not result in sheared arcades associated with filament channels.  Instead, evolving a MF model over multiple rotations was needed to recover the skew of coronal fields observed above the polarity inversion lines along which filaments appear.  Agreement between observations of filament handedness and their modeled counterparts was found to improve as the model was run for longer (multiple months) periods of time~\\citep*{Yeates:2007,Yeates:2008,Yeates:2009a,Yeates:2010}, indicating that the build-up and transport of energy over these time scales is likely an important factor.  The same evolving MF scheme was used in studies investigating the processes and frequencies by which filaments lift off~\\citep{MackayVanBallegooijen:2006a, MackayVanBallegooijen:2006b,Yeates:2009b,Su:2011}, the results of which showing that the changing magnetic geometry associated with neighboring bipolar active regions is critical in determining how the upward ejection of flux ropes occurs.  Until now, detailed NLFFF models of active region coronae have primarily been constructed only from single, isolated vector magnetograms, and therefore incorporate no information about the prior evolution of the photospheric and coronal magnetic field.  With the advent of the Helioseismic and Magnetic Imager~\\citep[HMI;][]{Scherrer:HMI,Schou:HMI} on board NASA's Solar Dynamics Observatory (SDO), time series of photospheric vector magnetogram data are now available at relatively high resolution ($0.5\"$ pixels) and cadence (every 12 min), with unprecedented spatial and temporal coverage. Such data enable detailed models of active regions to be constructed, wherein the MF scheme is utilized to advance the models in time in response to driving from time series of photospheric magnetograms.  This technique provides a way to accommodate the (possibly widespread) instances where the prior history of the state of the coronal magnetic field factors into the determination of its future state.  We show in a series of articles that it is feasible to use a MF approach with a time-evolving lower boundary condition to model the coronal evolution of an active region over a week-long period of time, and that ejections of magnetic flux can be driven in a model corona using this approach.  In this article, the first in the series, we outline the numerical methods and demonstrate the scheme on NOAA Active Region (AR) 11158 using idealized boundary conditions.  AR 11158, which was on disk in mid Feb.~2011 and produced several major flares, is followed for a significant fraction of its disk passage. ",
        "conclusions": "In this paper, we present a framework for performing data-driven simulations of solar AR formation. Under the magnetofriction assumption, fluid velocities are assumed to be proportional to the local Lorentz force. This has the advantage that the velocity distribution is guaranteed to evolve the field toward a force-free state. Since the magnetic induction equation is solved to advance the magnetic field in time, \\emph{ideal} magnetofrictional relaxation preserves the topology of the magnetic field while decreasing magnetic energy. In regions of high current density where resistive diffusion (used in this Eulerian code to facilitate magnetic reconnection) is important, the magnetic topology is allowed to change via reconnection. The change in the topology may permit the magnetic configuration to further relax to lower energy states under quasi-ideal, magnetofrictional evolution.  By incorporating time sequences of photospheric magnetograms as boundary data, the method models changes in the coronal magnetic field in response to photospheric driving (including shearing flows and flux emergence). As an example of the application of this method to magnetogram data obtained by SDO/HMI, we performed a number of simulations to model NOAA AR 11158. Since the accurate retrieval of photospheric electric fields from temporal sequences of vector magnetograms is still a research problem~\\citep{Fisher:EstimatingElectricFields}, we choose to perform the simulations with varying assumptions about the underlying $\\vec{E}$ that is responsible for the magnetogram evolution. In the case when twisting motion is absent, flux rope ejections were produced by the model. When continuous twisting was imposed, the sheared field above the sharp polarity inversion line in the core of the AR produced a series of flux rope ejections.  Since idealized assumptions were made about the nature of the underlying photospheric electric field, the simulations presented in this paper are not meant to be faithful representations of actual eruptions from  AR 11158. Nevertheless, this work demonstrates the potential utility of such a data-driven approach for modeling observed ARs. In a follow up paper, we will drive the model with actual electric fields retrieved from HMI vector magnetograms and constrained by observed Doppler velocities~\\citep{Fisher:EfieldsWithDoppler}. The use of electric fields more faithful to the data will facilitate a side-by-side comparison of the modeled AR with EUV observations from AIA~\\citep{Lemen:AIA}.  A method for synthesis of mock `coronal' images base on a proxy emissivity was introduced. The emissivity of points along a magnetic field line is assumed to proportional to the field-line average of $j^2$. This simple technique seems to produce with a visual texture similar to EUV images of coronal loops. While this technique is useful for visualizing an ensemble of coronal loops in a simple magnetic model, it is by no means a replacement for more sophisticated techniques that use the thermodynamic variables from MHD models and take into account atomic physics for EUV image synthesis."
    },
    "1208/1208.0951_arXiv.txt": {
        "abstract": "{ Using detailed Monte Carlo simulations we have characterized the features of the radio emission of inclined air showers in the Ultra-High Frequency band (300 MHz -  3 GHz). The Fourier-spectrum of the radiation is shown to have a sizable intensity well into the GHz frequency range. The emission is mainly due to transverse currents induced by the geomagnetic field and to the excess charge produced by the Askaryan effect. At these frequencies only a significantly reduced volume of the shower around the axis contributes coherently to the signal observed on the ground. The size of the coherently emitting volume depends on frequency, shower geometry and observer position, and is interpreted in terms of the relative time delays with respect to a plane shower front moving at the speed of light. At ground level, the maximum emission at high frequencies is concentrated in an elliptical ring-like region around the intersection of a Cherenkov cone with its vertex at shower maximum and the ground.  The frequency spectrum of inclined showers when observed at positions that view shower maximum in the Cherenkov direction, is shown to be in broad agreement with the pulses detected by the Antarctic Impulsive Transient Antenna (ANITA) experiment, making the interpretation that they are due to Ultra-High Energy Cosmic Ray atmospheric showers consistent with our simulations. These results are also of great importance for experiments aiming to detect molecular bremsstrahlung radiation in the GHz range as they present an important background for its detection. } ",
        "introduction": "Radio pulses from ultra-high energy cosmic rays (UHECRs) were first observed in the 1960s and 1970s in coincidence with air shower arrays~\\cite{jelley,allan} but severe difficulties were encountered to extract sufficiently precise information about the showers themselves using this technique alone~\\cite{Fegan_review}. Current advances in fast signal digitization and increased computational power have opened up new possibilities to exploit the technique and multiple efforts are being carried out to characterize these pulses in detail. Experiments such as AERA~\\cite{AERA} in the context of the Pierre Auger Collaboration, LOPES~\\cite{LOPES} in the context of KASCADE, and LOFAR~\\cite{LOFAR}, are regularly detecting pulses. These experiments are typically operating in the $\\sim$ 30-80 MHz range since the emission is expected to be at least partially coherent at those frequencies. There are also efforts to detect GHz emission from EAS claimed to originate from a different and yet unconfirmed mechanism: molecular bremsstrahlung~\\cite{Gorham_molecular}, such as AMBER~\\cite{Gorham_molecular}, MIDAS~\\cite{MIDAS}, CHROME~\\cite{CHROME} and EASIER~\\cite{microwave-auger}.  Recently several events in the 300-900~MHz frequency range were serendipitously observed with the ANITA balloon-borne antenna array which was flown over Antarctica.  These events have been claimed to be consistent with emission from Extended Air Showers (EAS) with energies between $10^{18}$~--~$5\\times10^{19}$~eV \\cite{ANITA_UHECR}. Earlier claims of detection of UHECRs in the ultra-high frequency (UHF) band exist \\cite{Fegan_Nature1,Spencer_Nature,Fegan_Nature2}. Other experiments have also recently observed GHz radiation associated with EAS~\\cite{EASIER,CHROME} although at higher frequencies, typically above 3 GHz. These results motivate a careful study of the radio emission induced by UHECR showers paying particular attention to the UHF band in the 300 MHz - 3 GHz range.  The emission is due to both shower charges and currents: an excess of electrons accumulates because there are no positrons in the atmosphere, a mechanism predicted by Askaryan~\\cite{Askaryan62}, and drift currents are induced through the separation of positively and negatively charged particles in the magnetic field of the Earth~\\cite{Kahn-Lerche,Allan_Nature,Falcke_radio_EAS,CODALEMA}, with the latter mechanism dominating the emission.  When the induced currents and the net charge distributions move at a speed faster than the speed of light in the atmosphere, several features characteristic of Cherenkov emission appear, including an enhancement of the high frequency components of the field \\cite{ZHAireS,Scholten_arXiv}. Calculations of the emission from air showers have been performed using Monte Carlo simulations~\\cite{ZHAireS,REAS3,SELFAS,COREAS}, and analytical and semi-analytical techniques~\\cite{Scholten_MGMR,Seckel_ARENA2012}.  In this work we characterize the features of the UHF emission from EAS using ZHAireS \\cite{ZHAireS}, a full Monte Carlo simulation of the air shower and its associated radio emission. The simulation is based on the AIRES shower code \\cite{aires}, and the emission is calculated using a well established algorithm \\cite{allan} obtained from first principles~\\cite{ZHS91,ZHS92,ARZ10} with which the electric field due to each individual charged particle track produced in the shower simulation is calculated. In this process interference between the emission from different tracks is carefully accounted for, including the variation of the refractive index $n$ with altitude. The algorithm does not assume any specific emission mechanism and correctly reproduces the radiation induced by the excess charge and the geomagnetic field \\cite{ZHAireS}. ",
        "conclusions": "The radio signal emitted from UHECR air showers has been shown to have a sizable intensity well into the GHz frequency range. In relative terms the UHF components are more important for inclined showers than for vertical ones. When projecting on the ground, the peak of the UHF signal is located at a narrow elliptical ring defined by the intersection of the Cherenkov cone at shower maximum and the ground. The width of the ring becomes narrower as the frequency is increased. The lateral distance to shower axis at which the UHF signal is recieved at highest amplitude depends on geometry, increasing with zenith angle, and reaching hundreds of meters for showers of $70^\\circ$.  We have shown that the lateral distribution for the low frequency components of the radio pulses, typically below $\\sim$100 MHz, displays a rather flat behavior with distance to the shower core until a position is reached at which shower maximum is viewed at the Cherenkov angle \\cite{ZHAireS,Scholten_PRL}. This behavior is consistent with observations made at LOPES~\\cite{LOPES_LDF} and LOFAR \\cite{LOFAR} and should be taken into account when trying to obtain the shower geometry and properties from the radio pulse measurements \\cite{AERA}.  The frequency spectrum of the pulses received in the Cherenkov cone has been shown to extend well into the GHz regime and to fall with an exponential slope in the UHF band, with a constant that is largest at the Cherenkov angle and drops as the observer views shower maximum away from the Cherenkov angle. Predictions for short radio pulses at ns scales follow from the spectra obtained for inclined showers.  In the case of more vertical showers, although the UHF components exist, the pulses are dominated by the lower frequency components and thus give typically a much broader pulse in time.  Our results indicate that by looking at the pulses at different distances from the shower core with a dense array it could be in principle possible to obtain information about the depth development of the shower \\cite{Scholten_arXiv}, but the detectors mut be positioned at fairly close distances because the Cherenkov cone is typically very narrow. These results provide an important background for experiments that are trying to measure molecular bremsstrahlung radiation from extensive air showers, and should be carefully considered in order to properly interpret measurements made in the UHF band. It is worth noting that molecular bremsstrahlung is thought to be isotropic \\cite{Gorham_molecular}, while the geomagnetic radiation studied in this paper is highly collimated in the direction of the Cherenkov cone.  The power spectra obtained in the Cherenkov direction are consistent with the spectra measured at ANITA and attributed to radio pulses from UHECR \\cite{ANITA_UHECR}. The results of the simulations made in this work further support this hypothesis. An event by event comparison of ANITA pulses and realistic predictions, followed by a simulation of the experiment would further strengthen this claim. The third flight of ANITA, scheduled for 2013-2014, is expected to detect hundreds of cosmic rays~\\cite{ANITA_UHECR}, and could provide a highly significant test of the expected UHECR pulse properties. The EVA (Exa-Volt Antenna) project \\cite{EVA} could increase the event rate by a factor of 10."
    },
    "1208/1208.5329_arXiv.txt": {
        "abstract": "We reproduce  the broadband  spectrum of  Crab pulsar, from UV to very high energy gamma-rays - nearly ten decades in energy, within the framework of  the cyclotron-self-Compton model. Emission is produced  by  two counter-streaming beams within the outer gaps, at distances above  $\\sim$ 20 NS radii. The outward moving beam  produces UV-$X$-ray photons via Doppler-booster cyclotron emission, and GeV photons  by Compton scattering the cyclotron photons produced by the inward going beam. The scattering occurs in the deep Klein-Nishina regime, whereby the  IC component provides a direct measurement of particle distribution within the magnetosphere.   The required plasma multiplicity is high, $\\sim 10^6-10^7$, but is consistent with the average particle flux injected into the pulsar wind nebula.  The importance of  Compton scattering in the Klein-Nishina regime also implies the importance of  pair production in the outer gaps. We suggest that outer gaps are important sources  of pairs  in pulsar magnetospheres.  Cyclotron motion of particles  in the pulsar magnetosphere may be excited due to  coherent  emission of radio waves by  streaming particles at the anomalous cyclotron resonance. Thus, a whole range of  Crab non-thermal emission, from coherent radio waves to very high energy $\\gamma$-rays - nearly eighteen  decades in energy - may be a manifestation of inter-dependent  radiation processes.  The present model, together with the observational evidence in favor of the IC scattering (Lyutikov et al. 2012; Lyutikov 2012), demonstrates that the inverse Compton scattering can be the dominant high energy emission mechanism in majority of pulsars. ",
        "introduction": "The pulsar high energy emission is a complicated  unsolved problem in high energy astrophysics. It has been  been under intensive study for nearly four decades \\citep{cr77,1986ApJ...300..500C,1995ApJ...438..314R,1996ApJ...458..278D,2008ApJ...680.1378H}. The   \\Fermi\\ Gamma-Ray Space Telescope detected  a large number of pulsars \\citep{2010ApJS..187..460A}; this  revolutionized our picture of the non-thermal emission from pulsars in the gamma-ray band from 100\\,MeV up to about 10\\,GeV. At even higher energies, in the very-high energy (VHE) band, the detection of the Crab pulsar at 25\\,GeV by the Magic Collaboration  \\citep{2008Sci...322.1221A} and recently at 120\\,GeV by the VERITAS Collaboration  \\citep{VERITASPSRDetection} in the very-high energy  band  allow to stringently constrain the very-high-energy emission mechanisms   {\\it Geometrical} models, based on the idea of the outer gap \\citep{1986ApJ...300..500C}, are very successful in explaining the basic features of the observed $\\gamma$-ray light curves \\cite[\\eg][]{1995ApJ...438..314R,2008ApJ...680.1378H,2010ApJ...715.1282B}. While there seems broad consensus that the particle accelerator is located in the outer magnetosphere, the radiation physics remain controversial \\citep[\\eg][]{1996A&AS..120C..49A}.    Nearly universally, the origin of the emission above $\\sim $ 100 MeV was until recently attributed to the curvature emission  \\citep[\\eg][]{cr77,1986ApJ...300..500C,1995ApJ...438..314R,1999MNRAS.308...54H,2003ApJ...591..334H,2007ApJ...662.1173H,2008ApJ...680.1378H}. For example, \\cite{1986ApJ...300..522C}  concluded that \"Crab primary outer gap $e^+/e^-$ lose most of their energy to curvature $\\gamma$-rays\". Curvature radiation has remained as the preferred gamma-ray emission mechanism \\citep{1996ApJ...470..469R} (see also \\cite{1986ApJ...300..500C,2000ApJ...537..964C,2008MNRAS.386..748T,2008ApJ...676..562T}). Possible importance of the IC scattering was discussed in application to the Vela pulsar \\citep[\\eg][]{1986ApJ...300..522C,1996ApJ...470..469R} \\citep[see also][]{1986ApJ...300..500C,2000ApJ...537..964C,2008ApJ...680.1378H}. The IC scattering was assumed to be done by the particles in the \\ms\\  interacting with surface $X$-ray photons. The IC scattering   was not deemed to be {\\it the dominant mechanism of high energy emission}.  In contrast, we argued   \\citep{2012ApJ...754...33L,2012arXiv1203.1860L} that  the  IC scattering may be the dominant source of high energy photons in a majority of pulsars \\citep[see also][]{2012A&A...540A..69A}. In this paper we further develop the IC model to include a modeling of the broadband SED, from UV to very high energy $\\gamma$-rays, covering nearly ten decades in energy. In its essence,  the lower energy UV-$X$-ray peak is due to the cyclotron emission by the secondary particles, Doppler boosted by the parallel motion of the plasma to the $X$-ray range, while the GeV component is due to the scattering of the cyclotron photons by the counter-streaming beam, Fig.  \\ref{picture}. ",
        "conclusions": "In this paper we demonstrate that the cyclotron-self-Compton model  is a viable model of pulsar high energy emission. Using   observationally-constrained  properties of the distribution function   of particles within the pulsar \\ms, as well as  fairly general additional theoretical assumptions, we are able to reproduce the overall SED of the Crab pulsar over nearly ten decades of energy. The key theoretical assumption is the presence of two counter-streaming populations within the outer gaps. The outward moving beam then produces UV-$X$-ray photons via Doppler-booster cyclotron emission, and the GeV emission by Compton scattering the cyclotron photons produced by the inward going beam. As a simplifying assumption, the parameters of the inward and outward going beams were chosen to be the same: if the distribution is skewed in favor of the outgoing beam, the intensity of the ingoing radiation will be reduced, producing a nearly one-sided emission pattern (in addition, cyclotron absorption may be important for the inward going radiation \\citep{1995JApA...16..173R,2000ApJ...537..964C}).   Naturally, the current one-dimensional model is a simplification; it is expected that more advanced models, that incorporate self-consistently the geometry, structure of accelerating electric fields and various radiative processes,  will result in the modification of the inferred parameters. The most important simplifications include,  first, the assumption of a IC scattering in a deep KN regime is not valid for scattering of the cyclotron photon emitted close to the \\LC\\ by the slowly moving particles. Second, the separation of the distribution function into parallel and perpendicular parts is likely not to be a unique from the observation point of view. For example, relativistic transverse motion vill increase the emitted frequency in the center-of-gyration frame, thus reducing the required parallel boost.   I would  like to thank Jonathan Arons, Charles Dermer, Kouichi Hirotani and George Machabeli for encouraging discussions."
    },
    "1208/1208.0210_arXiv.txt": {
        "abstract": "Tensor modes in the cosmic microwave background are one of the most robust signatures of inflation. We derive theoretical bounds on the tensor fraction, as a generalization of the well-known Lyth bound. Under reasonable assumptions, the new bounds are at least two orders of magnitude stronger than the original one. We comment on a previously derived generalization, the so-called  Efstathiou-Mack relationship. We also derive a new absolute upper bound on tensors using de Sitter entropy bounds. ",
        "introduction": "Understanding the initial conditions that led to structure formation in our Universe is one of the most important issues in modern cosmology. Among the variety of available scenarios, inflation occupies a special place. In addition to being a theoretically attractive paradigm, it is compatible with all the available observational data.  In its simplest version, single field slow-roll inflation is realized through a  canonically normalized scalar field minimally coupled to Einstein gravity. To match with observations, the height and the slope of the potential have to obey special relationships. However, in most of the models, the excursion of the scalar field (inflaton) is at least of the order of the reduced Planck scale\\footnote{Throughout the paper, we are using natural units $\\hbar=c=1$.} $M_P=(8 \\pi G_N)^{-1/2}$. Such large excursions potentially undermine the validity of effective field theory  used to derive the predictions. This is especially important considering the fact that quantum gravity corrections are likely to spoil the delicate balance between the height and the slope of the potential once the inflaton is allowed to travel over Planckian distances in field space\\footnote{Notice that, in addition to spoiling the flatness of the potential, non-renormalizable operators $\\lambda_n \\phi^{n+4}/M_P^n$ (with $n\\ge 1$) will make the energy density $\\gg M_P^4$ once the inflaton takes super-Planckian values.}. Inflation, in this sense, is unique because it is doubly UV sensitive: first because of the mass of the inflaton owing to its scalar nature, and second because of the  unavoidable proliferation of dangerous UV-suppressed operators. Of course, supersymmetry addresses the first issue, however the fine-tuning of the infinite tower of Planck-suppressed operators inevitably remains.  On the other hand, inflation predicts a stochastic background of gravitational waves (GWs) whose magnitude is related to the energy scale  of inflation \\cite{Lyth:1984yz} and, more importantly, to the inflaton  excursion  \\cite{Lyth:1996im}. According to the {\\em Lyth bound} \\cite{Lyth:1996im}, positive detection of tensors  would mean super-Planckian values of the inflaton which is clearly interesting both from the model building standpoint and from the observational one.  This  would also imply serious conceptual rethinking about effective field theory. Given the fact  that inflation offers an unequal and unique opportunity to access the Planck scale, and in the absence of experimental guidance on that question, it is important to use theoretical consistency to make some progress. In this article, we scrutinize this issue and derive theoretical bounds on the tensor fraction. \\vspace{-4mm} ",
        "conclusions": "In conclusion, we reiterate the truism that the tensor fraction is a very important observable to test the initial conditions of the Universe. We derived two theoretical bounds on that quantity. The first bound Eq.~(\\ref{boundr}) means that (for sub-Planckian inflaton excursion and thus consistent field theory description) $r\\lesssim 0.002$, placing it beyond the reach of {\\sl Planck} but within reach of {\\sl COrE} and {\\sl PIXIE} \\cite{Kogut:2011xw}. On the other hand, the fact that both single-field benchmark scenarios that are consistent with current data have  $\\Delta\\phi \\gtrsim 10\\, M_P$ make their tensor fraction  within reach of {\\sl CMBPol}. The second bound,  arising from de Sitter entropy bounds,  implies an upper bound on tensors Eq.~(\\ref{boundent}), independently of the magnitude of inflaton excursion, which is within reach of {\\sl Planck} and future observations. The detection of $B$-modes would promote inflation from an attractive paradigm to a predictive theory. However this is not the end of the story, as any realization of slow-roll inflation will have to face the issue of a consistent UV-completion.  We end up with some speculations. It would be interesting to explore the relationship between the entropy of the inflaton and $r$. For instance, in \\cite{Conlon:2012tz} it was argued that the entropy of the inflaton is proportional to $(\\Delta\\phi/H)^2$, which by the de Sitter entropy bound,  would censor any attempt to have super-Planckian inflaton excursion, and thus observable inflationary GWs, in a consistent UV-complete theory. This is in agreement with the conclusions reached in \\cite{Banks:2003sx}, where it was shown that attempts to build natural inflation models in string theory with a decay constant $f\\gg M_P$ are doomed to failure. \\vspace*{-5mm}"
    },
    "1208/1208.3025_arXiv.txt": {
        "abstract": "We use cosmological simulations to study the effects of self-interacting dark matter (SIDM) on the density profiles and substructure counts of dark matter halos from the scales of spiral galaxies to galaxy clusters, focusing explicitly on models with cross sections over dark matter particle mass $\\sigma/m = 1$ and $0.1$ $\\hbox{cm}^2/\\hbox{g}$.  Our simulations rely on a new SIDM N-body algorithm that is derived self-consistently from the Boltzmann equation and that reproduces analytic expectations in controlled numerical experiments.   We find that well-resolved SIDM halos have constant-density cores, with significantly lower central densities than their CDM counterparts.   In contrast, the subhalo content of SIDM halos is only modestly reduced compared to CDM,   with the suppression greatest for large hosts and small halo-centric distances. Moreover, the large-scale clustering and halo circular velocity functions in SIDM are effectively identical to CDM, meaning that all of the large-scale successes of CDM are equally well matched by SIDM. From our largest cross section runs we are able to extract scaling relations for core sizes and central densities over a range of halo sizes and find a strong correlation between the core radius of an SIDM halo and the NFW scale radius of its CDM counterpart. We construct a simple analytic model, based on CDM scaling relations, that captures all aspects of the scaling relations for SIDM halos.  Our results show that halo core densities in $\\sigma/m = 1 \\, \\hbox{cm}^2/\\hbox{g}$ models are too low to match observations of galaxy clusters, low surface brightness spirals (LSBs), and dwarf spheroidal galaxies.  However, SIDM with  $\\sigma/m \\simeq 0.1 \\, \\hbox{cm}^2/\\hbox{g}$ appears capable of reproducing reported core sizes and central densities of dwarfs, LSBs, and galaxy clusters without the need for velocity dependence. Higher resolution simulations over a wider range of masses will be required to confirm this expectation. We discuss constraints arising from the Bullet cluster observations, measurements of dark matter density on small-scales and subhalo survival requirements, and show that SIDM models with $\\sigma/m \\simeq 0.1 \\, \\hbox{cm}^2/\\hbox{g} \\simeq 0.2\\, \\hbox{barn}/\\hbox{GeV}$ are consistent with all observational constraints. ",
        "introduction": "There is significant evidence that some form of dark matter dominates the gravitating mass in the universe and its abundance is known to great precision \\citep{komatsu11}.   The most popular candidate for dark matter is the class of weakly interacting massive particles (WIMPs), of which supersymmetric neutralinos are examples \\citep{steigman1985,griest1988,jungman1996}.  WIMPs are stable, with negligible self-interactions, and are non-relativistic at decoupling (``cold``).  It is important to recognize that of these characteristics, it is primarily their coldness that is well tested via its association with significant small-scale power.   Indeed, WIMPs are the canonical Cold Dark Matter (CDM) candidate.  Cosmological models based on CDM reproduce the spatial clustering of galaxies on large scales quite well \\citep{reidetal10} and even the clustering of galaxies on $ \\sim 1$ Mpc scales appears to match that expected for CDM {\\em subhalos} \\citep{kravtsov2004,conroy2006,trujilloGomezetal11,reddick2012}.  Beyond the fact that the universe appears to behave as expected for CDM on large scales,  we have few constraints on the microphysical parameters of the dark matter, especially those that would manifest themselves at the high densities associated with cores of galaxy halos.     It is worth asking what (if anything) about vanilla CDM can change without violating observational bounds.  In this paper we use cosmological simulations to explore the observational consequences of a CDM particle that is strongly self-interacting, focusing specifically on the limiting case of velocity-independent, elastic scattering.  Dark matter particles with appreciable self-interactions have been discussed in the literature for more than two decades \\citep{carlson92,machacek93,delaix95,spergelandsteinhardt00,firmani00}, and are now recognized as generic consequences of hidden-sector extensions to the Standard Model \\citep{pospelov08, arkanihamed09, ackerman09, feng09, feng10a, loeb11}. Importantly, even if dark sector particles have no couplings to Standard Model particles they might experience strong interactions with themselves, mediated by dark gauge bosons (see \\citealt{feng10} and \\citealt{peter12} for reviews).   {\\em The implication is that  astrophysical constraints associated with the small-scale clustering of dark matter may be the only way to test these scenarios}.  Phenomenologically, self-interacting dark matter (SIDM) is attractive because it offers a means to lower the central densities of galaxies without destroying the successes of CDM on large scales. Cosmological simulations that contain only CDM indicate that dark-matter halos should be cuspy and with (high) concentrations that correlate with the collapse time of the halo \\citep{nfw97, bullock2001, wechsler2002}.  This is inconsistent with observations of galaxy rotation curves, which show that galaxies are less concentrated and less cuspy than predicted in CDM simulations \\citep[e.g.][]{floresandprimack94,simonetal05, kuzioetal08, bloketal10, dutton2011, kuzioetal11,ohetal11a, walkerandpenarrubia11, saluccietal12, castignanietal12}.  Even for clusters of galaxies, the density profiles of the host dark-matter halos appear in a number of cases to be shallower than predicted by CDM-only structure simulations, with the total (dark matter + baryons) density profile in a closer match to the CDM prediction for the dark matter alone \\citep[e.g.][]{sandetal04, sand2008, newmanetal09, newmanetal11,coe2012,umetsu2012}.  One possible answer is feedback.   In principle, the expulsion of gas from galaxies can result in lower dark matter densities compared to dissipationless simulations, and thus bring CDM models in line with observations \\citep{governato2010,ohetal11b,pontzen2011,brooketal12,governato2012}.  However, a new level of concern exists for dwarf spheroidal galaxies \\citep{mbketal11a, ferreroetal11,mbketal11b}.  Systems with $M_* \\sim 10^6 \\Msun$  appear to be missing $\\sim 5 \\times 10^7 \\Msun$ of dark matter compared to standard CDM expectations  \\citep{mbketal11b}.  It is difficult to understand how feedback from such a tiny amount of star formation could have possibly blown out enough gas to reduce the densities of dwarf spheroidal galaxies to the level required to match observations (\\citealt{mbketal11b,teyssieretal2012,zolotov2012,penarrubia2012}; Garrison-Kimmel et al., in preparation).  \\cite{spergelandsteinhardt00} were the first to discuss SIDM in the context of the central density problem (see also \\citealt{firmani00}). The centers of SIDM halos are expected to have constant density isothermal cores that arise as kinetic energy is transmitted from the hot outer halo inward \\citep{balberg2002,colinetal02,ahn2005,koda2011}. This can happen if the cross section over mass of the dark matter particle, $\\sigmam$, is large enough for there to be a relatively high probability of scattering over a time $t_{\\rm age}$ comparable to the age of the halo:  $\\Gamma \\, t_{\\rm age} \\sim 1$, where $\\Gamma$ is the scattering rate per particle. The rate will vary with local dark matter density $\\rho(r)$ as a function of radius $r$ in a dark halo as \\begin{equation}\\label{eq:gamma} \\Gamma(r) \\simeq \\rho(r) (\\sigma/m) v_{\\mathrm{rms}}(r) \\, , \\end{equation} up to order unity factors, where $v_{\\mathrm{rms}}$ is the rms speed of dark-matter particles.   Based on rough analytic arguments, \\citet{spergelandsteinhardt00} suggested  $\\sigmam \\sim 0.1-100 \\ \\cmspg$ would produce observable consequences in the cores of halos.  Numerical simulations have confirmed the expected phenomenology of core formation \\citep{bukert00} though \\citet{kochanek&white00} emphasized the possibility that SIDM halos could eventually become {\\em more} dense than their CDM counterparts as a result of eventual heat flux from the inside out (much like core collapse globular clusters).  However this process is suppressed when merging from hierarchical formation is included \\citep[for a discussion see][]{ahn2005}.  We do not see clear signatures of core collapse in the halos we analyzed for $\\sigmam=1 \\ \\cmspg$.  The first cosmological simulations aimed at understanding dwarf densities were performed by \\citet{daveetal01} who used a small volume ($4 \\hMpc$ on a side) in order to focus computational power on dwarf halos.  They concluded that $\\sigmam = 0.1-10 \\cmspg$ came close to reproducing core densities of small galaxies, favoring the upper end of that range but not being able to rule out the lower end due to resolution.  Almost concurrently, \\citet{yoshida00} ran cosmological simulations focusing on the cluster-mass regime.  Based on the estimated core size of cluster CL 0024+1654, they concluded that cross sections no larger than $\\sim 0.1 \\ \\cmspg$ were allowed, raising doubts that constant-cross-section SIDM models could be consistent with observations of both dwarf galaxies and clusters.  These concerns were echoed by \\citet{miralda2002} who suggested that SIDM halos would be significantly more spherical than observed for galaxy clusters. Similarly, \\citet{gnedinandostriker01} argued that SIDM would lead to excessive sub halo evaporation in galaxy clusters.  More recently, the merging cluster system known as the Bullet Cluster has been used to derive the limits (68\\% C.L.) $\\sigmam < 0.7  \\cmspg$ \\citep{randalletal08} based on evaporation of dark matter from the subcluster and $\\sigmam <1.25  \\cmspg$  \\citep{randalletal08} based on the observed lack of offset between the bullet subcluster mass peak and the galaxy light centroid. In order to relax this apparent tension between what was required to match dwarf densities and the observed properties of galaxy clusters, velocity dependent cross sections that diminish the effects of self-interaction in cluster environments have been considered \\citep{firmani00,colinetal02,feng09,loeb11,vogelsberger12}.  There are a few new developments that motivate us to revisit constant SIDM cross sections on the order of $\\sigmam \\sim 1 \\cmspg$.    For example, the cluster (CL 0024+1654) used by \\citet{yoshida00} to place one of the tightest limits at $\\sigmam = 0.1$,  is now recognized as an ongoing merger along the line of sight \\citep{czoske2001,czoske2002,zhang2005,jee2007,jee2010,umetsu2010}.  This calls into question its usefulness as a comparison case for non-merging cluster simulations.  In a companion paper (Peter, Rocha, Bullock and Kaplinghat, 2012) we use the same simulations described here to show that published constraints on SIDM based on halo shape comparisons are significantly weaker than previously believed. Further, the results presented below clearly demonstrate that the tendency for subhalos to evaporate in SIDM models \\citep{gnedinandostriker01} is not significant for $\\sigmam \\sim 1 \\cmspg$. Finally (and related to the previous point), the best numerical analysis  of the Bullet Cluster \\citep{randalletal08} used initial cluster density profiles that were unmotivated cosmologically with central densities about a factor of two too high for the SIDM cross sections considered (producing a scattering rate that is inconsistently high).  Based on this observation, the bullet cluster constraint based on evaporation of dark matter from the subcluster should be relaxed since the amount of subcluster mass that becomes unbound is directly proportional to the density of dark matter encountered in its orbit. Moreover, their model galaxies were placed in the cluster halo potentials without subhalos surrounding them, an assumption (based on analytic estimates for SIDM subhalo evaporation) that is not supported by our simulations. This could affect the constraints based on the (lack of) offset between dynamical mass and light. Thus we believe that the bullet cluster constraints as discussed above are likely only relevant for models with $\\sigmam > 1 \\cmspg$.  However, the constraints could be made significantly stronger by comparing SIDM predictions to the densities inferred from the convergence maps since the central halo densities for $\\sigmam \\simeq 1 \\cmspg$ are significantly lower than the CDM predictions, as we show later.  Given these motivations, we perform a set of cosmological simulations with both CDM and SIDM.  For SIDM we ran $\\sigmam = 1$ and $0.1 \\cmspg$ models (hereafter SIDM$_1$ and SIDM$_{0.1}$), {\\em i.e.}, models that we have argued pass the Bullet cluster tests. Our simulations provide us with a sample of halos that span a mass range much larger than either \\citet{daveetal01} or \\citet{yoshida00} both with and without self-interactions.  One of the key findings from our simulations is that the core sizes are expected to scale approximately as a fixed fraction of the NFW scale radius the halo would have in the absence of scatterings. We can see where this scaling arises from a quick look at Equation~\\ref{eq:gamma}. This equation allows us to argue that the radius ($r_1$) below which we expect dark matter particles (on average) to have scattered once or more is set by: \\begin{equation}\\label{eq:onescatter} \\rhos f(r/\\rs) v_\\mathrm{rms} \\propto \\frac{\\Vmax^2}{\\Rmax^2} f(r_1/\\rs) \\Vmax = \\rm constant \\enspace, \\end{equation} where $f(x)$ is the functional form of the NFW (or a related) density profile. In writing the above equation we have assumed that the density profile for SIDM is not significantly different from CDM at $r_1$, something that we verify through our simulations. Now, since CDM enforces a $\\Vmax-\\Rmax$ relation such that $\\Vmax\\propto \\Rmax^{1.4-1.5}$, we see that the solution to $r_1/\\rs$ is going to be only mildly dependent on the halo properties. We develop an analytic model based on this insight later, but this is the underlying reason for why we find core sizes to be a fixed fraction of the NFW scale radius of the same halo in the absence of scatterings.  The major conclusion we reach based on the simulations and the analytic model presented here is that a self-interacting dark matter model with a cross-section over dark matter particle mass $\\sim 0.1 \\cmspg$ would be capable of reproducing the core sizes and central densities observed in dark matter halos {\\em at all scales}, from clusters to dwarf spheroidals, without the need for velocity-dependence in the cross-section.  In the next section, we discuss our new algorithm to compute the self-interaction probability for N-body particles, derived self-consistently from the Boltzmann equation. We discuss this new algorithm in detail in Appendix \\ref{appendixA}. In \\S \\ref{implementation.sec}, we show how this algorithm is implemented in the publicly available code GADGET-2 \\citep{springel05}. We run tests that show that our algorithm gets the correct interaction rate and post-scattering kinematics. The results of these tests are in \\S \\ref{test.sec}. The cosmological simulations with this new algorithm are described in detail in \\S \\ref{sims.sec}. In \\S \\ref{prelim.sec} we provide some preliminary illustrations of our simulation snapshots and in \\ref{lss:sec} we demonstrate that the large-scale statistical properties of SIDM are identical to CDM.   In \\S \\ref{halos.sec} we present the properties of individual SIDM$_1$ and SIDM$_{0.1}$ halos and compare them to the their CDM counterparts. In \\S \\ref{subhalos.sec} we discuss the subhalo mass functions in our SIDM and CDM simulations and show that SIDM$_1$ subhalo mass functions are very close to that of CDM in the range of halo masses we can resolve. We provide scaling relations for the SIDM$_1$ halo properties in \\S \\ref{scaling.sec} and in \\S \\ref{analytic.sec} we present an analytic model that reproduces these scaling relations as well as the absolute densities and core radii of SIDM$_1$ halos. We use these scaling relations and the analytic model to make a broad-brush comparison to observed data in \\S \\ref{discussion.sec}. We present a summary together with our final conclusions in \\S \\ref{sumandconc.sec}. ",
        "conclusions": "\\label{sumandconc.sec}  We have presented a new algorithm to include elastic self-scattering of dark matter particles in N-body codes and used it to study the structure of self-interacting dark matter (SIDM) halos simulated in a full cosmological context.  Our suite of simulations (summarized in Table 1) rely on identical initial conditions to explore SIDM models with velocity-independent cross sections $\\sigma/m = 1 \\ \\cmspg$ and $\\sigma/m = 0.1 \\  \\cmspg$  as well as a comparison set of standard CDM simulations (with $\\sigma/m = 0$).  Our primary conclusion is that while SIDM looks identical to CDM on large scales, SIDM halos have constant density cores, with core radii that scale in proportion to the standard CDM scale radius ($r_{\\rm core} \\simeq \\epsilon \\, \\rs$).    The relative size of the core increases with increasing cross section ($\\epsilon \\simeq 0.7$ for $\\sigma/m = 1$ and $\\epsilon \\sim 0.2$ for $\\sigma/m = 0.1 \\ \\cmspg$). Correspondingly, at fixed halo mass, core densities decrease with increasing SIDM cross section.  For both core radii and core densities, there is significant scatter about the scaling with $\\Vmax$ of the halo. The scaling relationship is strong enough that measurements of dark matter densities in the cores of dark matter dominated galaxies and large galaxy clusters likely provide the most robust constraints on the dark matter cross section at this time.  In a companion paper (Peter, Rocha, Bullock and Kaplinghat, 2012) we demonstrate, contrary to previous claims, that SIDM constraints from halo shape measurements may be less restrictive than (or at least similar to those from) measurements of absolute core densities alone.  Based on our simulation results we conclude that the dark matter self-scattering cross section must be smaller than $1 \\ \\cmspg$ in order to avoid under-predicting the observed core densities in galaxy clusters, low surface brightness spirals (LSBs), and dwarf spheroidal galaxies.  However, an SIDM model with a {\\em velocity-independent} cross section of about $\\sigma/m = 0.1 \\ \\cmspg$ appears capable of reproducing reported core sizes and central densities of dwarfs, LSBs, and galaxy clusters.   Higher resolution simulations with better statistics will be needed to confirm this expectation.  \\bigskip  \\noindent An accounting of our results are as follows:  \\begin{itemize} \\item Outside of the central regions of dark matter halos ($r \\gtrsim 0.5 R_{\\rm vir}$)  the large scale properties of SIDM cosmological simulations are effectively identical to CDM simulations.  This implies that all of the large-scale confirmations of the CDM theory apply to SIDM as well. \\\\  \\item The subhalo $\\Vmax$ function in SIDM with $\\sigma/m = 1 \\ \\cmspg$ differs by less than $\\sim 30\\%$ compared to CDM across the mass range $5\\times 10^{11}M_\\odot - 2\\times 10^{14}M_\\odot$ studied directly with our simulations . Differences in the $\\Vmax$ function with respect to CDM are only apparent deep within the centers of large dark-matter halos.  Thus, although is possible, it will be difficult to constrain SIDM models based on the effects subhalo evaporation.  \\\\  \\item SIDM produces halos with constant density cores, with correspondingly lower central densities than CDM halos of the same mass. For $\\sigmam = 1 \\ \\cmspg$, our simulated halo density structure is reasonably well characterized by a Burkert (1995) profile fit with a core size $r_{\\rm b} \\simeq 0.7 \\rs$, where $\\rs$ is the NFW scale radius of the same halo in the absence of self-interactions.  Core densities tend to increase with decreasing halo mass ($\\rho_b \\propto \\Mvir^{-0.2}$) but demonstrate about a factor of $\\sim 3$ scatter at fixed mass (likely owing to the intrinsic scatter in dark matter halo concentrations). \\\\  \\item SIDM halo core sizes, central densities, and associated scaling relations can be understood in the context of a simple analytic model. The model treats the SIDM halo as consisting of a core region, where self-interactions have redistributed kinetic energy to create an approximately isothermal cored density profile; and an outer region, where self-interactions are not effective. The transition between these regions is set by the strength of the self-interactions and this model allows us to make quantitative predictions for smaller cross sections where the cores are not resolved by our simulations.  Based on this model and a few of our best resolved simulated halos we find core sizes $\\sim 0.1 \\rs$ for $\\sigmam = 0.1 \\ \\cmspg$. \\\\  \\item Halo core densities over the mass range from $10^{15} - 10^{10} \\Msun$ in SIDM with $\\sigmam = 1 \\ \\cmspg$  are too low  ($ \\sim 0.005 - 0.04 \\, \\Msun/\\pc^3$) to match observed central densities in galaxy clusters ($\\sim 0.03 \\Msun/\\pc$) and dwarf spheroidals ($\\sim 0.1 \\Msun/\\pc^3$).  \\\\   \\item Halo core central densities in SIDM with $\\sigmam = 0.1 \\ \\cmspg$ are in line with those observed from galaxy clusters to tiny dwarfs  ($0.02 - 0.5 \\Msun/\\pc^3$) without the need for any velocity dependence.  The densities are more consistent with observations than those predicted in dissipationless CDM simulations, which are generically too high. SIDM models with this cross section over dark matter particle mass value are consistent with Bullet cluster observations, subhalo survival requirements and, as we show in a companion paper (Peter, Rocha, Bullock and Kaplinghat, 2012), measurements of dark matter halo shapes.  \\end{itemize}   Future work is necessary to expand both the dynamic range of our simulations in halo mass and resolution as well as the dynamic range in cross sections.  These simulations are necessary in order to make detailed comparisons with observations given the exciting possibility that dark matter self-interaction with $\\sigmam$ in the ballpark of $0.1 \\ \\cmspg$ could be an excellent fit to the central densities of halos over 4-5 orders of magnitude in mass."
    },
    "1208/1208.4433_arXiv.txt": {
        "abstract": "We introduce a model of {\\it potential driven DBI Galileon inflation}  in background ${\\cal N}$=1, ${\\cal D}$=4 SUGRA. Starting from D4-$\\bar{D4}$ brane-antibrane in the bulk ${\\cal N}$=2, ${\\cal D}$=5 SUGRA including quadratic Gauss-Bonnet corrections, we derive an effective ${\\cal N}$=1, ${\\cal D}$=4 SUGRA by dimensional reduction, that results in a Coleman-Weinberg type Galileon potential. We employ this potential in modeling inflation and in subsequent study of primordial quantum fluctuations for scalar and tensor modes. Further, we estimate the major observable parameters in both {\\it de Sitter (DS)} and {\\it beyond de Sitter (BDS)} limits and confront them with recent observational data from WMAP7 by using the publicly available code CAMB. ",
        "introduction": "In the very recent days a good number of theoretical physicists have devoted their attention to the development of consistent modified gravity theories which play analogous role as dark energy or the cosmological constant \\cite{ratra}, \\cite{caroll}. In higher-dimensional setups as in the case of {\\it DGP model} \\cite{rdvali},\\cite{def} where self-acceleration is sourced by a scalar field (the zero helicity mode of the 5D graviton), these types of Infra-Red (IR) modification of gravity \\cite{nima}, \\cite{gia} play a crucial role. Moreover, the  {\\it  DGP model} replicates the general relativistic features due to non-linear interactions via the well known {\\it Vainshtein mechanism} \\cite{vain,suji,babi}. Despite its profound success it has got some serious limitations \\cite{suji}, \\cite{royma}, which are resolved by introducing a dynamical field, {\\it aka}, Galileon \\cite{espo,nicolis,felice} arising on the brane from the bulk in the DGP setup. The cosmological consequences of the Galileon models have been studied to some extent in \\cite{fabio,chow,kobayashi,ajtoll}. Very recently, a natural extension to the scenario has been brought forward by tagging Galileon with the good old DBI model \\cite{guss}, \\cite{edmu}, resulting in so-called ``DBI Galileon'' \\cite{claudia,trodden,goon,shun,sebar}, that has reflected a rich structure from four dimensional cosmological point of view. However, in most of the physical situations, this type of {\\it effective} gravity theories are plagued with additional degrees of freedom which often results in unwanted debris like ghosts, Laplacian instabilities etc \\cite{de,dec1,feli}. Very recently a nobel effort towards the supersymmetric extension of the DBI Galileon model and its inflationary signature are discussed in \\cite{sazi} and \\cite{gaziql}. In the present paper we introduce a single scalar field model described by the D3 DBI Galileon originated from D4-$\\bar{D4}$ brane anti-brane setup in the background of five dimensional local version of the supersymmetric theory (supergravity). This prevents the framework from having extra degrees of freedom as well as {\\it Ostrogradski instabilities} \\cite{ost},  resulting in a higher-order derivative scalar field theory free from any such unwanted instabilities. Nevertheless, a consistent field theoretic derivation of the effective potential commonly used in the context of DBI Galileon cosmology has not been brought forth so far. On top of that, it is imperative to point out that the SUGRA origin of D3 DBI Galileon is yet to be addressed. In this article we plan to address both of these issues explicitly by deriving the inflaton potential from our proposed framework of DBI Galileon. It turns out that the derived action includes, in certain limits, the decoupling limit of DGP model as well as some consistent theories of massive gravity; and it also includes the ``K-mouflage'' \\cite{zio}, \\cite{zioa} and also ``G/KGB'' inflation \\cite{kobayashi}, \\cite{yama,kamada,vik}. Moreover, in general appearance of non-vanishing frame functions in the 4D action expedites breakdown of shift symmetry. Without shift symmetry, it may happen that the theory is unstable against large renormalization. The background action chosen in our model preserves shift symmetry of the single scalar field which gives it a firm footing from phenomenological point of view as well.  The plan of the paper is as follows. First we propose a fairly general framework by taking the full DBI action in D4 brane in the background of bulk ${\\cal N}$=2, ${\\cal D}$=5 supergravity \\cite{sayan1,sayan2,sayan6,nills} including the quadratic modification in Einstein's Hilbert action via Gauss-Bonnet correction in the bulk coming from two loop correction in string theory \\cite{sayan19}. Hence, using dimensional reduction technique, we derive the effective action for DBI Galileon in D3 brane in the background of ${\\cal N}$=1, ${\\cal D}$=4 supergravity induced by the  quadratic correction in the geometry sector and study cosmological inflationary scenario therefrom. We next engage ourselves in studying quantum fluctuations, by employing second order perturbative action for scalar and tensor modes in {\\it de Sitter (DS)} and {\\it beyond de Sitter (BDS)} limits, and hence calculate the primordial power spectrum of the scalar and tensor modes, their running and other observable parameters. We further confront our model with observations by using the publicly available code CAMB \\cite{camb}, and find them to fit well with latest observational data from WMAP7\\cite{wmap7} and expected to fit fair well with upcoming data from PLANCK\\cite{planck}. ",
        "conclusions": ""
    },
    "1208/1208.6093_arXiv.txt": {
        "abstract": "{Dynamo action in giant planets and rapidly rotating stars leads to a broad variety of magnetic field geometries including small scale multipolar and large scale dipole-dominated topologies. Previous dynamo models suggest that solutions become multipolar once inertia becomes influential. Being tailored for terrestrial planets, most of these models neglected the background density stratification.} {We investigate the influence of the density stratification on convection-driven dynamo models.} {Three-dimensional nonlinear simulations of rapidly rotating spherical shells are employed using the anelastic approximation to incorporate density stratification. A systematic parametric study for various density stratifications and Rayleigh numbers at two different aspect ratios allows to explore the dependence of the magnetic field topology on these parameters.} {Anelastic dynamo models tend to produce a broad range of magnetic field geometries that fall on two distinct branches with either strong dipole-dominated or weak multipolar fields. As long as inertia is weak, both branches can coexist but the dipolar branch vanishes once inertia becomes influential. The dipolar branch also vanishes for stronger density stratifications. The reason is the concentration of the convective columns in a narrow region close to the outer boundary equator, a configuration that favors non-axisymmetric solutions. In multipolar solutions, zonal flows can become significant and participate in the toroidal field generation. Parker dynamo waves may then play an important role close to onset of dynamo action leading to a cyclic magnetic field behavior.} {These results are compatible with the magnetic field of gas planets that are likely generated in their deeper conducting envelopes where the density stratification is only mild. Our simulations also suggest that the fact that late M dwarfs have dipolar or multipolar magnetic fields can be explained in two ways. They may differ either by the relative influence of inertia or fall into the regime where both types of solutions coexist. } ",
        "introduction": "The magnetic fields of planets and rapidly rotating low-mass stars are maintained by magnetohydrodynamic dynamos operating in their interiors. Scaling laws for convection-driven dynamos successfully predict the magnetic field strengths for both types of objects indicating that similar mechanisms are at work \\citep{Christensen09}. Recent observations show a broad variety of magnetic field geometries, ranging from large-scale dipole-dominated topologies \\citep[as on Earth, Jupiter and some rapidly rotating M dwarfs, see e.g.][]{Donati06,midM} to small-scale more complex magnetic structures \\citep[e.g.][]{earlyM,lateM}. Explaining the different field geometries remains a prime goal of dynamo theory. Unfortunately, the extreme parameters of planetary and stellar dynamo regions can not directly be adopted in the numerical models so that scaling laws are of prime importance \\citep{Christensen10}.  Global simulations of rapidly rotating convection successfully reproduce many properties of planetary dynamos \\citep[e.g.][]{Christensen06}. They show that the ordering influence of the Coriolis force is responsible for producing a dominant large scale dipolar magnetic field \\citep[for a stellar application, see also][]{Brown10} provided the Rayleigh number is not too large. Multipolar geometries are assumed once the Rayleigh number is increased beyond a value where inertial forces become important \\citep{Kutzner02,Sreenivasan06}. \\cite{Christensen06} suggest that the ratio of (nonlinear) inertial to Coriolis forces can be quantified with what they call the ``local Rossby number'' $\\text{Ro}_\\ell=u_{\\text{rms}}/\\Omega \\ell$, where $\\ell$ is the typical lengthscale of the convective flow. Independently of the other system parameters, the transition between dipole-dominated and multipolar magnetic fields always happens around $\\text{Ro}_\\ell \\simeq 0.1$. Scaling laws then, for example, predict a dipole-dominated field for Jupiter and a multipolar field for Mercury \\citep{Olson06}. The latter may explain the strong quadrupolar component in the planet's magnetic field \\citep{Christensen06nat}. Earth lies at the boundary where the field is dipolar most of the time but occasionally forays into the multipolar regime allowing for magnetic field reversals.  The scaling laws based on Boussinesq simulations are geared to model the dynamo in Earth's liquid metal core where the background density and temperature variations can be neglected. Their application to gas giants and stars where both quantities vary by orders of magnitude is therefore questionable \\citep{Chabrier06,Nettelmann12}. Recent compressible dynamo models of fully-convective stars suggest a strong influence of the density stratification on the geometry of the magnetic field: while the fully compressible models of \\cite{Dobler06} have a significant dipolar component, the strongly stratified anelastic models of \\cite{Browning08} tend to produce multipolar dynamos \\citep[see also][]{Bessolaz11}. These differences stress the need of more systematic parameter studies to clarify the influence of density stratification.  In addition, most of the previous Boussinesq studies have employed rigid flow boundary conditions for modelling terrestrial dynamo models. Stress-free boundary conditions, more appropriate for gas planets and stars, show a richer dynamical behaviour, including hemispherical dynamos and bistability where dipole-dominated and multipolar dynamos coexist at the same parameters \\citep[e.g.][]{Grote00,Busse06,Goudard08,Simitev09,Sasaki11}. This challenges the usefulness of the $\\text{Ro}_\\ell$ criterion since bistable cases exist significantly below $\\text{Ro}_\\ell\\simeq 0.1$ \\citep[e.g.][]{Schrinner12}.  Here we adopt the anelastic approximation \\citep[e.g.][]{Glatz1,Brag95,Lantz99} to explore the effect of a background density stratification on the dynamo process while filtering out fast acoustic waves. We conduct an extensive parameter study where we vary the degree of stratification and the Rayleigh number to determine the parameter range where dipole-dominated fields can be expected. The anelastic approximation and the numerical methods are introduced in section~\\ref{sec:model}. In section~\\ref{sec:results}, we present the results of the parametric study and describe the different dynamo regimes before relating our results to observations in the concluding section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have investigated the influence of background density stratification on convection-driven dynamos in a rotating spherical shell. The use of the anelastic approximation allowed us to exclude sound waves and the related short time steps \\citep{Glatz1,Clune99,Jones09}. Previous Boussinesq results have shown that inertial effects play a decisive role for determining the magnetic field geometry. When inertia becomes influential, only multipolar solutions with weak magnetic fields are possible. When inertia is weak, two solutions can coexist: dipole-dominated solutions with strong magnetic fields are found for stress-free as well as rigid boundary conditions. For stress-free and mixed boundary conditions, a multipolar branch is found at identical parameters \\citep[e.g.][]{Simitev09,Schrinner12}. Alternatively, the recent study by \\cite{King12} suggests that the transition between dipolar and multipolar dynamos may occur when inertia becomes larger than viscous forces. This transition is accompanied by an abrupt decrease of the kinetic helicity. These results however seem  to be in contradiction with previous studies where inertia is always larger than viscosity \\citep[e.g.][]{Wicht10} and further investigations are required to clarify this contradiction.  Our anelastic simulations confirm this scenario for mild stratifications corresponding to $N_\\rho<1.8$. The reason for the bistability is a competition between zonal winds and dipolar magnetic fields. Strong dipolar magnetic fields  prevent significant zonal winds to develop. Strong zonal winds, on the other hand prevent the production of significant dipolar fields. The two branches also differ in the induction mechanism. Strong zonal winds promote an $\\Omega$-effect which leads to a $\\alpha\\Omega$ or  $\\alpha^2\\Omega$ type of dynamo, while dipole-dominated magnetic fields are typically generated in an $\\alpha^2$ process. The sizable axisymmetric toroidal fields produced by an $\\Omega$-effect typically lead to a coherent cyclic time evolution of the magnetic field for moderate magnetic Reynolds numbers ($\\text{Rm} < 200$). This is consistent with previous Boussinesq studies \\citep[e.g.][]{Schrinner07} and numerical models of young solar-type stars \\citep{Brown11} and has been identified as Parker waves. Contrary to what is observed in the solar cycle, these waves start at the equator and travel towards the poles because of the opposite sign in the zonal shear.  For stronger stratification with $N_\\rho>1.8$, the dipolar branch is lost and close to onset of dynamo action a new magnetic mode characterised by a large wave number $m=2$ appears. This is likely due to the concentration of the $\\alpha$-effect into a narrow region close the equator. According to mean-field models, this would preferentially promote non-axisymmetric dynamos \\citep[see also][]{Chabrier06} with large wave numbers. The collapse of the dipolar branch may explain the differences between the weakly stratified and significantly dipolar simulations by \\cite{Dobler06} ($\\rho_{\\text{bot}}/\\rho_{\\text{top}} \\simeq 5$, i.e.~$N_\\rho=1.6$) and the strongly stratified anelastic and multipolar models by  \\cite{Browning08} ($\\rho_{\\text{bot}}/\\rho_{\\text{top}} \\simeq 100$, i.e.~$N_\\rho=4.6$).  Numerical limitations force us to use excessively large diffusivities in our simulations. Ekman numbers are thus orders of magnitude too large and Reynolds number orders of magnitude too low. When extrapolated to fast rotating planets and stars this type of models nevertheless provides, for example, realistic magnetic field strengths \\citep{Olson06,Christensen09}. This gives us confidence to compare the observed field geometries of planets and stars with predictions based on our simulation results.  Our result are compatible with the dipole-dominated magnetic fields on Jupiter and Saturn that are generated in their deeper metallic envelopes where the density stratification is only mild \\citep[roughly $N_\\rho\\sim 1.5-2$, see][]{Heimpel11,Nettelmann12}. Since the local Rossby number is very small for these planets, however, bistability seems to be an option. This is also the case for Uranus and Neptune, where $\\text{Ro}_\\ell$ and the density contrast within the dynamo region are small \\citep{Hubbard91,Olson06}. Their multipolar magnetic field would then suggest that these planets would then occupy the alternative branch offered by the bistability phenomenon.  Concerning rapidly rotating low-mass stars that may also fall into the low $\\text{Ro}_\\ell$ regime, the spectropolarimetric observations of \\cite{lateM} suggest that late M stars with very similar parameters (mass and rotation rate) come in two categories: some stars present a strong dipole-dominated magnetic field while others show weaker and multipolar magnetic structures. These two geometries may represent the two coexisting dynamo branches at smaller local Rossby numbers.  For the bistability to be a viable explanation, however, our simulations suggest that the dynamos must operate in a region with moderate density stratification ($N_\\rho \\sim 1-2$) and supercriticality ($\\text{Ra}/\\text{Ra}_c \\sim 10-30$). Since the stellar dynamos operate presumably far from onset of convection, only multipolar fields would then be possible. However, when decreasing the Ekman number towards more realistic values, the simulations by \\cite{Christensen06} suggest that the dipolar window may persist at higher supercriticalities. In addition, since these stars have huge density contrasts, our investigation would then generally predict multipolar fields. A further exploration of the parameter space seems here required to clarify this point. \\cite{Stanley09}, for example, suggest that lower Prandtl number may help in creating stronger dipole fields. Considering radial-dependent properties (e.g. viscosity, thermal diffusivity and electrical diffusivity) is also known to have a strong impact on the location of the convective columns that could possibly help to avoid the concentration of helicity close to the equator."
    },
    "1208/1208.1399_arXiv.txt": {
        "abstract": "{ We present a method for beam-deconvolving cosmic microwave background (CMB) anisotropy measurements. The code takes as input the time-ordered data along with the corresponding detector pointings and known beam shapes, and produces as output the harmonic $a_{Tlm}$, $a_{Elm}$, and $a_{Blm}$ coefficients of the observed sky. From these one can derive temperature and Q and U polarisation maps. The method is applicable to absolute CMB measurements with wide sky coverage, and is independent of the scanning strategy. We tested the code with extensive simulations, mimicking the resolution and data volume of Planck 30GHz and 70GHz channels, but with exaggerated beam asymmetry. We applied it to multipoles up to $l=1700$ and examined the results in both pixel space and harmonic space. We also tested the method in presence of white noise. } ",
        "introduction": "Removing systematic effects plays an important role in the data analysis of modern high-sensitivity cosmic microwave background (CMB) experiments, such as the WMAP and Planck missions \\citep{jarosik-etal-2011,tauber-etal-2010}. In this work we concentrate on one source of systematic effects, the beam asymmetry. We present an efficient beam deconvolution code, called \\artdeco. It can be applied to any absolute CMB experiment with sufficient sky coverage.  The required input consists of the time-ordered data (TOD), the corresponding detector pointings, and known beam shapes. The primary output of the code consists of the harmonic $a_{Tlm}$, $a_{Elm}$, and $a_{Blm}$ coefficients of the sky. From these one can derive temperature and Q and U polarisation maps.  This is the usual deconvolution map-making problem \\citep[see, e.g.,][] {armitage-wandelt-2004,armitage-wandelt-2009,harrison-etal-2011}. Our method differs from the earlier works by an efficient reformulation of the deconvolution problem through heavy use of Wigner functions. The formulation is general, and makes no assumptions on the scanning pattern. Similar ideas have been studied by G.~Pr\\'ezeau, independently of us (private communication).  In the CMB literature the concept of map-making often refers to a procedure that involves removal of correlated noise. Methods for doing this have been treated in several papers \\citep{keihanen-etal-2005,keihanen-etal-2010,poutanen-etal-2006,ashdown-etal-2007b, ashdown-etal-2007a,kurki-suonio-etal-2009,sutton-etal-2009}. The methods discussed in these works construct pixelised temperature and Q, U polarisation maps, but do not attempt to correct for an asymmetric beam shape.  Our method steps in at the point where the correlated noise has already been cleaned from the data. The input TOD is assumed to include signal and uncorrelated noise only. We are assuming here that the correlated noise component can be removed from the data to a sufficient degree prior to the beam deconvolution step. The noise removal step itself is affected by beam asymmetries because signal differences caused by beam mismatch can be falsely interpreted as noise. Fortunately, in destriping methods it is possible to strongly reduce this effect by masking out the regions with the strongest foreground contamination that generate most of the effect. Also, destriping methods produce as natural output the noise-cleaned TOD, which is the required input in deconvolution map-making. Accordingly, destriping methods are the natural counterparts of our deconvolution method.  Since the aim of this paper is to present a new deconvolution method, rather than assess the performance of any particular instrument, we tested the code with a somewhat idealised simulation. In particular, we ignored other systematic effects such as frequency response. On the other hand, we chose strongly asymmetric beam shapes to show the beam effects more clearly. ",
        "conclusions": "We have presented an efficient beam deconvolution code, designed for absolute CMB experiments, and tested it on simulated data. The code is released under the terms of the GNU General Public License and can be obtained from \\url{http://sourceforge.net/projects/art-deco/}.  We looked at maps constructed from the recovered $a_{Tlm}$, $a_{Elm}$, and $a_{Blm}$ coefficients, and examined the coefficients directly in harmonic space. These reflect different aspects of the solution. We compared the results to a harmonic expansion of a binned map. We have also shown that with sufficient computational resources, we can extend the method to \\lmax=1700, which would be sufficient for the Planck 70GHz channel.  In absence of noise we could recover the $a_{Tlm}$ coefficients to a high accuracy, and remove the effects of beam asymmetry. When white noise was added, we were able to reach a lower value of \\lmax, but the advantage of deconvolution over binning was still clear.  In the case of polarisation, deconvolution worked well in absence of noise. We were able to almost completely remove the temperature leakage due to beam mismatch. When noise was added, results were less clear. Deconvolution removed the visible galactic residual that arises from beam mismatch, but in harmonic space deconvolution did not seem to bring a clear benefit over binning. The recovered $a_{Elm}$ coefficients were dominated by noise at nearly all multipoles, but this was true for both deconvolved and binned maps.  Our simulation used beams that were highly asymmetrical and also had a strong mismatch between them, which leads to a significant polarisation leakage. We did this to show the beam effects more clearly, but at the same time we also made the deconvolution problem very challenging. The accuracy we can expect when the code is applied to a real experiment strongly depends on the beam shapes of the particular experiment.  Some topics for future study can be mentioned. We have set a very strict convergence criterion for the CG iteration. Our convergence study indicates that a much more relaxed criterion could be sufficient for practical purposes. A future improvement would be to define a more suitable convergence criterion. This could reduce the computation time by a significant factor.  We applied the method to a simulation data set that provides full sky coverage. Our code is built on a formalism that makes no assumptions on the sky coverage. In practice, good convergence requires nearly full coverage. This is intuitively evident, since we are solving the harmonic coefficients, which necessarily represent the complete celestial sphere. If the input data only cover a part of the sky, the coefficients cannot be well constrained. A straighforward extension of the method would be to insert a prior that constrains the solution in the region which is not covered by the data, but this is beyond the scope of this paper and is a subject for future study.  Yet another topic for further study is determining the noise bias present in the TT and EE spectra at high multipoles.  Though the simulations used in this work mimic some aspects of the Planck mission, the parameters used do not reflect the properties of the actual instrument. The results presented in this paper are thus not representative of the sensitivity of the Planck experiment, and the authors do not represent the Planck collaboration in this context.  \\appendix"
    },
    "1208/1208.0333_arXiv.txt": {
        "abstract": "We describe the design, construction, and expected performance of two new fiber integral field units (IFUs) --- HexPak and GradPak --- for the WIYN 3.5m Telescope Nasmyth focus and Bench Spectrograph. These are the first IFUs to provide formatted fiber integral field spectroscopy with simultaneous sampling of varying angular scales. HexPak and GradPak are in a single cable with a dual-head design, permitting easy switching between the two different IFU heads on the telescope without changing the spectrograph feed: the two heads feed a variable-width double-slit. Each IFU head is comprised of a fixed arrangement of fibers with a range of fiber diameters. The layout and diameters of the fibers within each array are scientifically-driven for observations of galaxies: HexPak is designed to observe face-on spiral or spheroidal galaxies while GradPak is optimized for edge-on studies of galaxy disks. HexPak is a hexagonal array of 2.9 arcsec fibers subtending a 40.9 arcsec diameter, with a high-resolution circular core of 0.94 arcsec fibers subtending 6 arcsec diameter. GradPak is a 39 by 55 arcsec rectangular array with rows of fibers of increasing diameter from angular scales of 1.9 arcsec to 5.6 arcsec across the array. The variable pitch of these IFU heads allows for adequate sampling of light profile gradients while maintaining the photon limit at different scales. ",
        "introduction": "\\label{sec:intro} HexPak and GradPak are two new formatted fiber optic integral field units (IFUs) for the WIYN 3.5m telescope Bench Spectrograph\\footnotemark. \\footnotetext{The WIYN Observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University, and the National Optical Astronomy Observatory.} These two IFUs are unique because they are the first formatted fiber IFUs to include multiple fiber diameters within the same fiber head. Including multiple fiber diameters allows each of these IFUs to simultaneously sample different angular scales within the same observation. The smaller fibers can gather light from higher surface brightness regions (e.g.\\ the core or midplane of a galaxy disk) while the larger fibers can collect light from fainter, more diffuse regions (e.g.\\ larger scale heights or scale radii of a disk), thereby enabling high S/N measurements to be obtained simultaneously at a range of spatial positions.   The two IFUs share the same cable and ``foot'' for mounting onto the WIYN Bench Spectrograph. Sharing the same cable minimizes the total volume required for routing within the WIYN telescope in an already over-filled system. Sharing the same foot results in a unique dual-slit design, allowing the two heads to be exchanged at the telescope without requiring changes to the spectrograph system. HexPak has a high-resolution core of fibers three times smaller in diameter than the surrounding fibers. As a hexagonal array with a circular, high-resolution core, HexPak is tailored for studies of radially-distributed, diffuse light sources, such as face-on galaxy disks, spheroidal galaxies, or star clusters. The GradPak head consists of five different fiber sizes, arranged in rows to form a gradient of fiber diameters from one edge of the array to the other. It is designed for integral field spectroscopy of edge-on galaxies, making it well-suited for studying extra-planar gas and stars in spiral galaxy disks.   Including multiple fiber diameters comes at a cost, however. In the case where the system spectral resolution is limited by the slit width, this will result in a varying spectral resolution, inversely proportional to the fiber diameter. In this case the maximum resolution will change by a factor of 3 for GradPak and 3.1 for HexPak, increasing from the largest to the smallest fibers. However, the smallest reimaged fiber sizes will have contributions from optical aberrations from the spectrograph. As a result, we expect the achievable resolution to increase only by a factor of 2--2.5 for HexPak.   The science impact of changing spectral resolution depends on the specific application. For example, the velocity dispersion of stars is expected to increase with scale height above the disk midplane in edge-on disk galaxies. For the study of stellar velocity dispersions in spheroidal or face-on disk galaxies, as another example, it would be advantageous to \\emph{increase} spectral resolution with radius, since systems become dynamically colder moving outward. This is opposite what these instruments deliver. On the other hand, at lower surface brightness the limits of S/N prevent useful information from being obtained at high spectral resolution, and in this sense these instruments provide a practical balance between signal and resolution. As we describe below, ample sky fibers are included for all fiber sizes to ensure excellent sky subtraction.  This instrument follows in the legacy of the excellent WIYN fiber IFUs DensePak\\cite{Barden98} and SparsePak\\cite{Bershady04,Bershady05}. The primary impetus behind this project was the increased throughput and image quality of the newly redesigned WIYN Bench Spectrograph\\cite{Barden94,Bershady08,Knezek10}. In the process an opportunity arose to rebuild the decommissioned DensePak IFU. The fiber from DensePak is being reused for the larger fibers in the new HexPak array, and most of the hardware in the cable ``foot'' housing that terminates the cable in the spectrograph room at WIYN is reused from the DensePak cable. HexPak contains additional fibers that were newly purchased for this project. GradPak is made entirely using new fiber. Additionally, all the cabling and head mount hardware is newly constructed.   In \\S\\ref{sec:design} we detail the key science drivers that served as a design target for the instrument, as well as describe the design challenges inherent in the design of formatted IFUs with multiple fiber diameters.  We detail the construction process in \\S\\ref{sec:construction} and summarize the project in \\S\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We are in the final construction phase of two new fiber optic IFUs, GradPak and HexPak. These IFUs will be the first formatted fiber IFUs to utilize multiple fiber diameters within the same IFU head. By including multiple fiber diameters these IFUs will greatly expand the spectroscopic capabilities of the WIYN 3.5m telescope, providing the ability to sample varying spatial scales within the same observation with the highly versatile Bench Spectrograph. This will enable observations simultaneously spanning a wide range in surface brightness to be optimized for the photon limit at spectral resolutions between 1000 $< \\lambda/\\Delta\\lambda <$ 30,000.   HexPak is designed to observe axi-symmetric surface brightness profiles with a $36\\arcsec\\times41\\arcsec$ hexagonal region sampled by 2.\\farcs8\\ fibers and a 6\\arcsec diameter high-resolution core sampled by 0.\\farcs94\\ fibers. GradPak is optimized for linear surface brightness gradients, with a stacked pseudo-slit design spanning $39\\arcsec\\times55\\arcsec$ using fibers ranging from 1.\\farcs9\\ to 5.\\farcs6\\ in diameter. Each of these IFUs presented unique challenges for successfully incorporating multiple fiber diameters within the same fiber head. We have described two methods for overcoming these challenges, one for radial arrangements of fibers and one for linear arrangements. Our solutions optimize the packing fraction of science fibers while also achieving regular and precise placement and configuration of fibers within each IFU."
    },
    "1208/1208.5516_arXiv.txt": {
        "abstract": "{} % {We aim at deriving the molecular abundances and temperatures of the hot molecular cores in the high-mass star-forming region NGC~6334I and consequently deriving their physical and astrochemical conditions.} {In the framework of the Herschel guaranteed time key program CHESS (Chemical Herschel Surveys of Star Forming Regions), NGC~6334I is investigated by using the Heterodyne Instrument for the Far-Infrared (HIFI) aboard the Herschel Space Observatory. A spectral line survey is carried out in the frequency range 480--1907~GHz, and further auxiliary interferometric data from the Submillimeter Array (SMA) in the 230~GHz band provide spatial information for disentangling the different physical components contributing to the HIFI spectrum. The spectral lines in the processed Herschel data are identified with the aid of former surveys and spectral line catalogs. The observed spectrum is then compared to a simulated synthetic spectrum, assuming local thermal equilibrium, and best fit parameters are derived using a model optimization package.} {A total of 46 molecules are identified, with 31 isotopologues, resulting in about 4300 emission and absorption lines. High-energy levels ($E_u>1000$~K) of the dominant emitter methanol and vibrationally excited HCN ($\\nu_2=1$) are detected. The number of unidentified lines remains low with 75, or $<$ 2\\% of the lines detected. The modeling suggests that several spectral features need two or more components to be fitted properly. Other components could be assigned to cold foreground clouds or to outflows, most visible in the \\ce{SiO} and \\ce{H2O} emission. A chemical variation between the two embedded hot cores is found, with more N-bearing molecules identified in SMA1 and O-bearing molecules in SMA2.} {Spectral line surveys give powerful insights into the study of the interstellar medium. Different molecules trace different physical conditions like the inner hot core, the envelope, the outflows or the cold foreground clouds. The derived molecular abundances provide further constraints for astrochemical models.} ",
        "introduction": "\\label{intro} The astrochemical composition of the interstellar gas depends on the physical environment and the evolutionary state. In the cold dense regions molecules are unsaturated, while in contrast saturated molecules are much more abundant in hot cores \\citep{herbst}. Unsaturated molecules are carbon rich with long carbon chains, like radicals as \\ce{C_n H} (n=2--8)  or cyanopolyynes \\ce{HC_n N} (n=3,5,7,9,11), whereas saturated organic molecules are hydrogen rich with single bonds of carbon like \\ce{CH3OCH3}. Molecular ions are predominantly unsaturated in the cold phase, because most hydrogenation reactions are endothermic or hindered by potential barriers \\citep{herbst2}. Ices on dust grains begin to form in the cold phase and are observed by broad absorption bands at mid-infrared wavelengths. Sequential hydrogenitation of oxygen or nitrogen forms water or ammonia, respectively. Embedded in the water-ices are relevant amounts of CO, and through further hydrogenation \\ce{CO -> HCO -> H2CO -> H2COH -> CH3OH}. Time-dependent astrochemical models suggest that abundances of certain molecules, such as \\ce{H2S} and \\ce{SO2}, might be used as an indicator of the evolutionary phase \\citep{Nomura,Herpin}, and greater chemical complexity can be expected for longer evolution timescales \\citep{Garrod}. Approximately 165 molecules have been detected in the ISM or circumstellar shells. \\footnote{A current list can be found at \\url{http://www.astro.uni-koeln.de/cdms/molecules}}  \\Citet{Dishoeckhifi} emphasized the importance of spectral line surveys for the research of star-forming regions. Line surveys give the possibility to obtain a census of all atoms and molecules and give insights into their thermal excitation conditions and dynamics by studying line intensities and profiles, which allows seprating different physical components. The discovery of new species or new excited lines from known molecules is possible. Furthermore, the wide range of energy levels and eventually opacities of one species allows a better determination of the temperature and abundance. The uniform recording with the same spectrometer in one observation reduces the calibration uncertainties.  The giant molecular cloud \\object{NGC~6334}, also known as the ``Cat's Paw Nebula'', lies in the Carina-Sagittarius Arm of the Milky Way at a distance of $(1.7 \\pm 0.3)$~kpc \\citep{distance}. Studies at infrared wavelengths revealed sites of massive  star formation \\citep{persi}. \\object{NGC~6334I} was first explored by \\citet{emerson} as the brightest infrared source in the northern part of the cloud, containing an OH maser and an ultracompact \\ion{H}{II} region. According to \\citet{emerson}, the pumping of the OH maser requires an energy source of a massive collapsing protostar ($M>30~M_{\\odot}$). \\ion{H}{II} regions are associated with high-mass stars (spectral classes O and B) where the UV radiation is strong enough to ionize hydrogen. \\citet{Russeil} studied the densest cores in NGC~6334 in molecular line and dust continuum emission. For NGC~6334I, they derive a total mass of $M$=1206~$M_{\\odot}$ from the continuum flux density at 1.2~mm, an average density of $n_{\\rm H_2}$=$1.2 \\times 10^6$~$\\rm cm^{-3}$ and an upper limit for the virial mass of $M_{\\rm vir}$=525~$M_{\\odot}$. The total mass estimate is relatively high, because a low value of the dust temperature of only 20~K was assumed, as compared to the value of T=100~K derived by \\citep{Sandell}, which leads to a factor of 6 difference of the mass estimate. While an average dust temperature of T=25~K is appropriate for most cores in the NGC~6334 complex \\citep{ngcdust}, the dust temperature is known to be much higher in NGC~6334I.  Interferometric studies revealed that NGC~6334I consists of several cores, labeled SMA1 to SMA4 \\citep{Hunter}. The separation distance between SMA1 and SMA2 is about 3.5$''$ or 6000~AU. SMA1 and SMA2 show rich spectra of molecular transitions and SMA3 coincides with the \\ion{H}{II} region which exhibits free-free and dust emission and is excited by the near-infrared source IRS 1E \\citep{buizer}. SMA4 only shows dust emission but no line emission, which is still unexplained. Table~\\ref{tNGC} summarizes characteristic values for NGC~6334I derived from the dust continuum emission at different wavelengths and, which classifiy the source as a low-luminosity high-mass hot core. The overall size of NGC~6334I is 10$''$ $\\times$ 8$''$, where 1$''$ equals to 1700~AU. \\begin{table} \\centering \\caption{Mass and temperature estimates for NGC~6334I} \\label{tNGC} \\begin{tabular}{cccccccc} \\hline \\hline Source & $M (M_{\\odot})$ & $T_{\\rm d}$ (K) & $\\beta$ & n ($\\rm cm^{-3}$) & $L_{\\rm bol}$ ($L_{\\odot}$) \\\\ \\hline Total\\tablefootmark{a} & 200 $\\pm$ 100 &   100    &   1.7 $\\pm$ 0.3  &  $1.2\\times 10^7$ &   $2.6 \\times 10^5$  \\\\ \\hline SMA1\\tablefootmark{b} & 17 $\\pm$ 50\\% &   100    &     &   &      \\\\ SMA2\\tablefootmark{b} & 11 $\\pm$ 50\\% &   100    &     &   &       \\\\ SMA3\\tablefootmark{b} & 33 $\\pm$ 50\\% &   60    &     &   &      \\\\ SMA4\\tablefootmark{b} & 36 $\\pm$ 50\\% &   33    &     &   &      \\\\ \\hline \\end{tabular} \\tablefoot{Values taken from \\tablefoottext{a} \\citet{Sandell} and \\tablefoottext{b} \\citet{Hunter}. $M$ is the mass, $T_d$ the average dust temperature, $\\beta$ the dust emissivity index, $n$ the hydrogen density and $L_{\\rm bol}$ the bolometric luminosity.} \\end{table} SMA1 and SMA2 are so-called hot molecular cores, where certain attributes apply according to \\citet{cesaroni}: a high temperature ($T\\ge$100~K), sizes smaller than 0.1~pc, large masses ($10-1000$~$M_{\\odot}$) and luminosities exceeding $10^4$~ $L_{\\odot}$. Furthermore they are often associated with water masers and ultracompact \\ion{H}{II} regions. Hot cores are heated by high-mass protostars which are embedded in a dusty gas envelope. Their rich molecular spectra originate from the evaporation of ice mantles covering the dust grains, which release at $T\\sim$100~K molecules to the gas phase. Molecules released from grain mantles then begin to drive a fast, high-temperature gas-phase chemistry, forming complex organic species like acetone, beginning from precursors like methanol. Their rotational, vibrational and radiative excitation leads to a characteristic emission spectrum. This explains why the spectrum is poorer in earlier stages like the colder protostellar core \\object{NGC~6334I(N)} \\citep{Brogan,Walsh}, 2$'$ or 1~pc north of NGC~6334I.  Multiple spectral line surveys have been carried out for NGC~6334I: \\citet{McCutcheon} 334--348~GHz, \\citet{Schilke} 459--461~GHz and 817--819~GHz, \\citet{Thorwirth} 88--115~GHz and 218--267~GHz, \\citet{Walsh} 84--116~GHz and the latest from \\citet{Kalinina} in the range 81--242~GHz. Besides, several maser lines have been observed from \\ce{NH3} \\citep{beuther}, \\ce{CH3OH} \\citep{walshm} and \\ce{H2O} \\citep{Migenes} as well as bipolar outflows in CO \\citep{outflow,Qiu} and SiO \\citep{sio}. This source is part of the CHESS program (Chemical Herschel Surveys of Star Forming Regions), a key program of the Herschel Space Observatory \\citep{Ceccarelli}. The aim is to study eight different sites and phases of star formation varying in mass, luminosity, evolutionary state, astrochemical composition etc. by conducting and comparing their spectra. In this article, the entire spectrum (500--1900~GHz) of NGC~6334I is analyzed which was observed by Herschel in 2010. Previously, some small sections of the spectrum were studied to find transitions of new molecules not observed before in the interstellar medium (ISM): Oxidaniumyl (\\ce{H2O+}) by \\citet{Ossenkopf} and chloronium (\\ce{H2Cl+}) by \\citet{Lis}. A detailed analysis of water has been made by \\citet{water} and of methylidyne (CH) by \\citet{Wiel}.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=270,width=1\\textwidth]{smaalln.pdf}} \\caption[smaall]{Integrated SMA line intensity map over all species. Overlaid in black contours is the continuum emission. The contour levels are from 90\\% to 10\\% of the maximum in steps of 10\\%. The large and small white circles represent the FWHM Herschel beam size for band 1 and band 7, respectively.} \\label{fsma} \\end{figure}  One difficulty is that if a large beam encompasses a complex, unresolved source structure, the spectrum can become complex due to different overlapping components. The spatial information from the SMA is complementary. The Herschel beam is larger than NGC~6334I so that the source is spatially unresolved, while the SMA beam can resolve the individual cores. The SMA maps thus reveal the source morphology and give insights into the distribution of various molecules. This helps to decompose the contributions from different components present in the HIFI spectrum and to differentiate between the two hot cores SMA1 and SMA2, the extended emission and outflows, see Fig.~\\ref{fsma}. \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics[trim=5cm 0cm 0cm 0cm,clip,angle=270,width=1\\textwidth]{allbands_small.pdf}} \\caption{Complete HIFI spectrum of NGC~6334I. B-bands are coloured in blue and bands 5--7 are smoothed by averaging over 5--10 channels.} \\label{allbands} \\end{figure*} ",
        "conclusions": "In this article, observational data from the HIFI instrument aboard the Herschel Space Observatory  and from the Submillimeter Array are analyzed for the high-mass star-forming region NGC~6334I. We employ a LTE radiative transfer model in order to get insights into the physical structure of the source. The molecular spectral line survey gives an overview of the chemical inventory of the molecular species and their excitation, whereas the high-resolution spectrometer allows tracing the kinematics of the different velocity components. The interferometer maps reveal the source  morphology and help to distinguish the embedded cores and resolve their sizes. The results fit well with the overall picture of a high-mass young stellar object and hot molecular cores described by \\citet{Dishoeck}.  Two main software tools, XCLASS and MAGIX, are used to demonstrate that modeling and handling of the vast HIFI data sets is possible. These tools allow to decompose the spectrum into contributions of individual molecules and their isotopologues. By assuming LTE condition and employing the LM fitting algorithm, a fast extraction of information is possible. In order to take line blends into account, a model for every molecule is produced, which is consistent over multiple HIFI bands. It should be clear that for other excitation conditions than LTE and for optically thick lines like CO this method reaches its limitations. Sophisticated 3-d radiative transfer codes such as RADMC-3D\\footnote{\\url{http://www.ita.uni-heidelberg.de/~dullemond/software/radmc-3d/}} and LIME \\citep{LIME} with the possibility of non-LTE calculations will be used in the future in order to model NGC~6334I into more detail, taking into account the geometry and non-homogeneous distribution of temperature and abundances. With the improved capabilities of ALMA (Atacama Large Millimeter Array), it will be possible to study the molecular chemistry and kinematics on sub-arcsecond scales ($<$~1000~AU) and to follow the fragmentation of protostellar cores with much improved mass sensitivity.  The molecular survey revealed a variety of chemical species, including saturated organic molecules. The astrochemical question is: how can these abundances be explained and how are they related to each other? The next step is therefore a careful comparison of the column densities presented here with predictions of astrochemical models and detailed comparisons with the molecular abundances and abundance ratios in other CHESS targets."
    },
    "1208/1208.0043_arXiv.txt": {
        "abstract": "We study supernova neutrino flux variations in the IceCube detector, using 3D models based on a simplified neutrino transport scheme. The hemispherically integrated neutrino emission shows significantly smaller variations compared with our previous study of 2D models, largely because of the reduced SASI activity in this set of 3D models which we interpret as a pessimistic extreme. For the studied cases, intrinsic flux variations up to about 100~Hz frequencies could still be detected in a supernova closer than about 2~kpc. ",
        "introduction": "Measuring the neutrino signal from a future galactic core-collapse supernova (SN) will provide valuable information about the conditions in the collapsing stellar core and in the forming compact remnant. In particular neutrinos and gravitational waves will serve as direct probes of the explosion mechanism. While neutrino energy deposition is most widely favored as the trigger and energy supply of the SN blast wave \\cite{Bethe:1985}, the success of this mechanism in iron-core SNe turns out to be tightly linked to the development of violent nonradial mass flows in the layer behind the stalled bounce shock (see Ref.~\\cite{Janka:2012} for a recent review).  Multi-dimensional hydrodynamical models have shown that nonradial gas motions naturally grow from small initial perturbations in the neutrino-heated, shocked accretion flow because of convective instability~\\cite{Bethe:1990} and the standing accretion shock instability (SASI)~\\cite{Blondin:2002sm}. Violent, quasi-periodic expansion and contraction (``sloshing'') of the shock can lead to considerable variation of the mass accretion rate to the proto-neutron star (PNS). In axi-symmetric (2D) models, where the sloshing motions are artificially constrained to proceed along the symmetry axis, these variations and the associated compression of the PNS ``surface'' layer can produce fluctuations of the observable (hemispherically integrated) neutrino emission of more than 10\\% in luminosity and of 0.5--1\\,MeV in the mean spectral energy; the polar emission variations are even larger~\\cite{Marek:2008qi}. Such large variations of the neutrino emission would be easily detectable for a galactic SN with IceCube or future megaton-class instruments, and the presence of Fourier components with frequencies of tens of Hz up to 200--300 Hz would provide important information on the SN core dynamics prior to the onset of the explosion \\cite{Lund:2010,Brandt:2011}. They could also allow one to constrain neutrino masses \\cite{Ellis:2012a} and to probe neutrino propagation over cosmic distances \\cite{Ellis:2012b}.  Because of the directing effect of the symmetry axis the SASI sloshing motions appear particularly strongly in 2D core-collapse simulations, independently of whether detailed neutrino transport \\cite{Marek:2008qi, Marek:2007gr, Mueller:2012jm, Mueller:2012jh, Brandt:2011, Suwa:2010}, simplified neutrino cooling and heating terms \\cite{Murphy:2008, Nordhaus:2010, Hanke:2011}, a schematic cooling prescription, or an outflow boundary at the bottom of the otherwise adiabatic accretion flow~\\cite{Blondin:2002sm, Blondin:2006} is used.  In three-dimensional (3D) models, low-order multipole SASI sloshing motions have so far been observed only with considerably reduced amplitudes and stochastically changing direction \\cite{Iwakami:2008} or no clear signatures of SASI sloshing activity were seen \\cite{Wongwathanarat:2010jm, Mueller:2012jw, Takiwaki:2012, Burrows:2012}. Indeed, it has been claimed that buoyancy-driven convection dominates post-shock turbulence \\cite{Murphy:2012db} and the SASI is ``at most a minor feature of SN dynamics'' \\cite{Burrows:2012}. Such conclusions, however, should be taken with caution. They rely on small sets of 3D models with idealized setups, e.g., only including simple neutrino heating and cooling terms but without detailed neutrino transport and its feedback on the PNS evolution. They therefore do not explore the important influence of PNS contraction driven by energy and lepton losses through neutrino emission as well as progenitor-specific differences, both of which were shown to be able to fundamentally change the dynamics of the postshock accretion flow \\cite{Marek:2007gr, Janka:2012, Mueller:2012jh}. Moreover, the conditions for developing SASI spiral modes \\cite{Blondin:2006,Iwakami:2009,Fernandez:2010} in SN cores are neither fully understood nor satisfactorily investigated. The growth timescale of such modes can be strongly reduced by even small amounts of rotation in the collapsing stellar matter \\cite{Yamasaki:2008} and their appearance is observed in experimental analogies \\cite{Foglizzo:2012}. Low-mode asymmetries, which are much bigger than seen in present, first 3D simulations and which impose quasi-periodic, large-amplitude variability on the PNS neutrino emission, can therefore not be excluded.  We here take the pessimistic point of view and apply our previous analysis of 2D models \\cite{Lund:2010} to several recent 3D explosion calculations which show only a low level of SASI activity and no signs of long-lasting sloshing motions \\cite{Mueller:2012jw}, in contrast to the 2D simulations of Marek, Janka and M\\\"uller~\\cite{Marek:2008qi} evaluated by Lund et al.~\\cite{Lund:2010}. Nevertheless, even for such unfavorable conditions our results suggest that the modulation of the neutrino emission below about 100~Hz will be detectable with IceCube for a galactic SN up to a distance of 2~kpc.  The outline of our paper is as follows. In Sec.~\\ref{sec:numerical} we summarize essential information about the simulated models and the expected signal rates in IceCube. In Sect.~\\ref{sec:power} we present the power spectra that can be deduced from a measurement and discuss their meaning. In Sect.~\\ref{sec:shotnoise} we explore the detectability of the intrinsic neutrino signal variations in competition with the shot noise of the IceCube measurement. Conclusions will be given in Sec.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}   We have shown that also for 3D SN models time variations of the neutrino emission lead to distinct peaks in the power spectra as previously seen for 2D results~\\cite{Lund:2010}. Despite considerably smaller amplitudes of the luminosity and mean energy fluctuations than in the 2D simulations, these features will still be detectable in the IceCube event rate for a future galactic SN within 2\\,kpc distance. Up to frequencies of about 150\\,Hz such features will hint to intrinsic flux variations, whereas above this frequency they will, with a high probability, be associated to statistical effects in the measurement of the signal.  A thorough discussion is hampered by the simplifications made in the 3D models and the small set of investigated cases. In particular, the neutrino signals evaluated in this work are based on an approximate grey treatment of the neutrino transport in models, in which the high-density core of the PNS was excised and replaced by a boundary condition. Nevertheless, the signals exhibit all familiar features also predicted by more sophisticated transport schemes in fully self-consistent 1D and 2D simulations. The largest emission variations are observed during the postbounce accretion phase until shortly after shock revival, when the SN core is stirred by violent hydrodynamical instabilities. Large-scale asymmetries and intermittent behavior characterize the accretion flows to the nascent NS during this phase. Interestingly, the power spectra of the 3D models show similar peaks in the same frequency range as those found in 2D, although the flow dynamics in 3D appears to be largely different. In particular, the strong SASI sloshing motions that are constrained to the symmetry axis in the 2D case and that are responsible for big emission variations, cannot be seen prominently in the 3D simulations.  If the existence of the frequency peaks is consolidated by better 3D models, these features will offer a powerful diagnostic for the dynamics in the SN core and the explosion mechanism. The role of turbulence and violent nonradial mass motions in the postshock layer during the accretion phase and on the way to shock revival is presently a matter of vivid debate \\cite{Burrows:2012,Murphy:2012db,Mueller:2012jh}. A better understanding of the characteristic imprints of the SASI and of convective overturn motions on the neutrino emission in 3D is therefore desirable.  Our previously investigated 2D models together with the 3D simulations analyzed in this work can be considered to span the range between optimistic and pessimistic cases with respect to the possible level of signal variability. It is not excluded that fully self-consistent 3D models will reveal higher fluctuation amplitudes of the neutrino emission. Stronger SASI activity and especially the development of spiral modes \\cite{Blondin:2007nat}, whose growth could be fostered by even a relatively small amount of rotation in the SN core, can not be excluded on grounds of the small set of presently available 3D simulations."
    },
    "1208/1208.5499_arXiv.txt": {
        "abstract": "We present high spatial resolution observations of Giant Molecular Clouds (GMCs) in the eastern part of the nearby spiral galaxy NGC 6946 obtained with the Combined Array for Research in Millimeter-wave Astronomy (CARMA).  We have observed $\\co$, $\\cotwo$ and $\\cother$, achieving spatial resolutions of $5\\farcs4\\times5\\farcs0$, $2\\farcs5\\times2\\farcs0$ and $5\\farcs6\\times5\\farcs4$ respectively over a region of 6 $\\times$ 6 kpc.  This region extends from 1.5 kpc to 8 kpc galactocentric radius, thus avoiding the intense star formation in the central kpc.  We have recovered short-spacing $u$-$v$ components by using single dish observations from the Nobeyama 45m and IRAM 30m telescopes.  Using the automated CPROPS algorithm we identified 45 CO cloud complexes in the $\\co$ map and 64 GMCs in the $\\cotwo$ maps.  The sizes, line widths, and luminosities of the GMCs are similar to values found in other extragalactic studies.  We have classified the clouds into on-arm and inter-arm clouds based on the stellar mass density traced by the 3.6 $\\mu$m map.  Clouds located on-arm present in general higher star formation rates than clouds located in inter-arm regions.  Although the star formation efficiency shows no systematic trend with galactocentric radius, some on-arm clouds -- which are more luminous and more massive compared to inter-arm GMCs -- are also forming stars more efficiently than the rest of the identified GMCs.  We find that these structures appear to be located in two specific regions in the spiral arms.  One of them shows a strong velocity gradient, suggesting that this region of high star formation efficiency may be the result of gas flow convergence. ",
        "introduction": "The formation and evolution of giant molecular clouds (GMCs), the sites where stars form, are thought to be influenced by the local properties of the environment and global galactic structure (\\citealt{2008MNRAS.385.1893D}; \\citealt{2009ApJ...700..358T}; \\citealt{2011ApJ...730...11T}; \\citealt{2009ApJ...699..850K}).  Although the internal physical properties of molecular clouds have been extensively studied (\\citealt{1987ApJ...319..730S}; henceforth S87), a more detailed understanding of their origin and evolution in different types of environments is needed.  Maps with resolution matched to the scale of individual clouds are the key to achieve this goal; however, such maps currently exist only for a few galaxies outside the Local Group.  Due to the coupled evolution between GMCs and their host galaxies, achieving a more complete understanding of molecular cloud formation and evolution processes represents a key goal for studies of galaxy evolution.  While several mechanisms have been proposed, the predominant mode for molecular cloud formation and evolution in spiral galaxies is not yet clear.  In terms of the arm structure, spiral galaxies can be classified as {\\it grand-design}, most of them consisting of two symmetrical spiral arms; {\\it multi-armed}, which show several non-symmetrical arms; or {\\it flocculent} galaxies with multiple shorter arms.  For grand-design galaxies, the spiral structure is generated by density waves induced by bars or companion galaxies.  The influence of the density waves in the ISM in these galaxies has been commonly modeled assuming a long lived, rigidly rotating spiral potential (\\citealt{2002ApJ...570..132K}; \\citealt{2004MNRAS.349..270W}; \\citealt{2006MNRAS.367..873D}).  On the other hand, flocculent and multi-armed galaxies are thought to develop from local gravitational instabilities, although flocculent galaxies can present weak density waves as well (\\citealt{1997ApJ...484..202T}).  In this case, numerical simulations model the galactic potential as time dependent structures, which can change on a timescale of $\\sim$100 Myr (\\citealt{2008MNRAS.385.1893D}; \\citealt{2011ApJ...735....1W}).  The formation of massive molecular clouds in grand-design galaxies is attributed to the passing of spiral density waves.  A particular feature of this type of galaxy is the formation of spurs perpendicular to the spiral arms due to the shearing of the molecular structures as they leave the arms (\\citealt{2002ApJ...570..132K}; \\citealt{2004MNRAS.349..270W}).  In their study of the northeastern spiral segment of IC 342, a grand design galaxy, \\citet{2010PASJ...62.1261H} showed that actively star forming GMCs are more massive, have smaller velocity dispersions and are more gravitationally bound than quiescent GMCs.  They interpreted the tendency for the star-forming GMCs to lie downstream with respect to the quiescent molecular clouds as evidence in favor of cloud a coalescence scenario, where massive molecular clouds ($\\sim 10^6\\ \\Msun$) are formed from coagulation of smaller clouds entering the spiral arms, and the star formation is ignited once the excess kinetic energy is dissipated through inelastic collisions of the clouds (Hirota et al. 2011).  In another grand design galaxy, M51, \\citet{2009ApJ...700L.132K} suggested that evolution of GMCs is driven by large scale galactic dynamics, where the most massive molecular complexes (Giant Molecular Associations, GMAs) are created by coagulation of GMCs in spiral arms.  As GMAs leave the arms, they suffer fragmentation by the strong shear motions, which would explain the spur structures observed extending from the arms into the inter-arm regions.  The remaining fragmented GMCs are not disassociated into atomic gas, and they become part of the molecular cloud population observed in inter-arm regions which survive until the passage of another spiral arm.  On the other hand, \\citet{2008MNRAS.385.1893D} presented a numerical simulation of flocculent galaxies.  They find that the gaseous spiral structure essentially traces the potential minimum, in contrast to the case of grand-design spirals where an offset between the spiral shock and the potential minimum is expected.  According to the simulations provided by \\citet{2008MNRAS.385.1893D} and \\citet{2006MNRAS.371..530C} the most massive structures are thought to be produced in regions where collision or merging between spiral arms can occur, yielding an enhanced star formation in those overdense regions, in contrast to the situation proposed for grand-design galaxies where the molecular clouds are formed by agglomeration within spiral arms.  Conversely, the formation of inter-arm molecular gas structures are proposed to be the result of local instabilities rather than fragmentation of GMA's, and the properties of the clouds are approximately similar across the disk, with no offsets between gas and star formation expected.  High resolution observations of flocculent galaxies have provided evidence in support of this picture.  For instance, \\citet{2003PASJ...55..605T} mapped $\\co$ over the southern arm of the flocculent galaxy NGC 5055.  They found no obvious offset between H$\\alpha$ and the molecular gas, and that on-arm and inter-arm clouds do not have significant differences in their properties.  Resolved GMCs in the Milky Way have been shown to be in approximate virial equilibrium obeying scaling relations known as Larson's laws (\\citealt{1981MNRAS.194..809L}).  These laws establish that velocity dispersion (or line width), size, and luminosity are correlated in the Milky Way GMCs.  The size-line width relation may reflect turbulent conditions in the molecular ISM.  The typical GMC temperature of 15--25 K produces thermal CO velocity dispersions of only $\\sim$0.1 $\\kms$, providing evidence of non-thermal supersonic turbulence within the clouds.  Nevertheless, whether the turbulence in the molecular ISM is internally or externally generated is still unclear (\\citealt{2007ARA&A..45..565M}; \\citealt{2011ApJ...730...11T}).  The other two Larson relations are the luminosity-line width and the luminosity-size relations.  In a recent work, \\citet{2008ApJ...686..948B} (henceforth B08) present a study of resolved properties of GMCs for a sample of dwarf galaxies along with two disk galaxies of the Local Group.  They find insignificant differences between Milky Way GMCs and GMCs in dwarf galaxies, with the latter following approximately the same Larson relations as Galactic clouds.  \\citet{2007prpl.conf...81B} arrived at a similar conclusion by analyzing observations of CO surveys for galaxies in the Local Group.  Due to sensitivity and resolution limitations, observations of resolved properties of GMCs in galaxies other than Milky Way are in an early stage (\\citealt{2012ApJ...744...42D}, hereafter DM12; \\citealt{2009ApJ...700L.132K}).  Only a few galaxies have been mapped in CO emission across the full extent of the optical disk with sufficient resolution to probe the $\\sim$100 pc size scales on which GMCs form.  High resolution CO surveys (\\citealt{1999ApJS..124..403S}; \\citealt{2003ApJS..145..259H}) have focused primarily on the central regions, whereas complete mapping using single-dish telescopes (\\citealt{2007PASJ...59..117K}; \\citealt{2009AJ....137.4670L}) has been obtained at resolutions of 11$\\arcsec$ to 15$\\arcsec$, corresponding to $\\sim$500 pc for nearby spirals, inadequate to resolve GMCs.  Only for M33 (BIMA survey) and the Magellanic Clouds (NANTEN surveys) has it been possible to conduct a fairly complete census of GMCs (\\citealt{2003ApJS..149..343E}; \\citealt{2008ApJS..178...56F}).  Using a combination of $\\co$ and $\\cotwo$ and pursuing a staged approach of starting with large-scale, low resolution maps and following up with high resolution imaging on smaller areas, CARMA makes it possible to study the overall distribution of GMCs, as well as the properties of individual GMCs, in galaxies outside the Local Group.  We have successfully used this observational strategy in the nearby spiral galaxy NGC 6946.  As indicated by the single-dish CO map from IRAM (\\citealt{2009AJ....137.4670L}), this galaxy is extremely rich in molecular gas, with CO emission extended to a large fraction of the optical radius.  Table \\ref{n6946-prop} shows the basic parameters of NGC 6946.  We have chosen NGC 6946 as our target because of its proximity, high CO surface brightness, low inclination, and the availability of high-quality datasets at a variety of wavelengths, including HI imaging from THINGS and multiband Spitzer imaging from SINGS (\\citealt{2003PASP..115..928K}).  The spiral structure in NGC 6946 appears complex, with K-band images revealing four prominent, asymmetric spiral arms (\\citealt{1995ApJ...452L..21R}).  While the central region of the galaxy has been well-studied in CO because of its high CO surface brightness (\\citealt{2007AJ....134.1827C}; \\citealt{2007PASJ...59..117K}; DM12), the eastern region has not been mapped at high resolution.  In this paper we present observations of the eastern region of NGC 6946.  Firstly, we made observations of $\\co$ using CARMA covering a region of size 6$\\times$6 kpc$^2$.  The angular resolution for the $\\co$ map was $\\sim$110 pc, which is not sufficient to resolve structures with size scales similar to GMCs (S87; B08).  Then, the second step was to make higher resolution $\\cotwo$ observations ($\\sim$ 50 pc) towards the brightest CO complexes to study the properties of individual GMCs.   We present our study as follows:  in Section \\ref{obs} we describe our observations of NGC 6946 using CARMA, and we describe the archival data at several wavelengths that we include in our analysis.  In Section \\ref{cprops} we summarize the technique used to identify GMCs and to measure their physical properties.  In Section \\ref{cl-prop} we present the scaling relations of the cloud properties, and we compare them with previous studies of Galactic and extragalactic clouds.  In Section \\ref{env-eff} we discuss whether the properties of the clouds differ between on-arm and inter-arm regions.  In Section \\ref{discuss} we discuss the implications of our results, and in Section \\ref{summary} we summarize the work presented in this paper.  \\begin{table} \\caption{Properties of NGC 6946.} \\centering \\begin{tabular}{lc} \\hline\\hline Morph. \\tablenotemark{a} & SABcd \\\\ R.A. (J2000)$^\\mathrm{a}$ & 20:34:52.3 \\\\ Decl. (J2000$^\\mathrm{a}$ & 60:09:14 \\\\ Distance (Mpc)$^\\mathrm{b}$ & 5.5 \\\\ Incl. (\\degrees)$^\\mathrm{c}$ & 33\\\\ P.A. (\\degrees)$^\\mathrm{c}$ & 243\\\\ \\hline \\multicolumn{2}{l}{{\\bf Notes.}} \\\\ \\multicolumn{2}{l}{$^\\mathrm{a}$ NASA/IPAC Extragalactic Database (NED).}\\\\ \\multicolumn{2}{l}{$^\\mathrm{b}$ \\citet{1988JBAA...98..316T}.}\\\\ \\multicolumn{2}{l}{$^\\mathrm{c}$  \\citet{2008AJ....136.2563W}.} \\end{tabular} \\label{n6946-prop} \\end{table} ",
        "conclusions": "\\label{discuss}  \\subsection{Molecular cloud properties} Despite the long-standing acceptance of the Larson scaling relations, recent studies have raised questions about the form and universality of the scaling relations. For instance, \\citet{2009ApJ...699.1092H}, reviewing the properties of the Galactic molecular clouds studied by S87, found LTE masses that are typically a factor of a few smaller than the virial masses derived by S87.  However, maybe a more remarkable finding is that the size-line width relation coefficient ($v_\\mathrm{o} \\equiv \\sigma_v/R^{1/2}$) is not universal, but depends on the mass surface density of the clouds.  \\citet{2009ApJ...699.1092H} attributed the variation in $v_\\mathrm{o}$ to differences in magnetic field strength among the clouds.  Alternatively, \\citet{2011MNRAS.411...65B} proposed that the dependence of $v_\\mathrm{o}$ on the mass surface density is consistent with molecular clouds in a process of hierarchical and chaotic gravitational collapse.  They suggested that, although hydrodynamic turbulence is required to induce the first dense condensations in the ISM, once the structures become bound, the gravity is the main driver of the internal motions.  Figure \\ref{figure_co_fact} shows the relation between $v_\\mathrm{o}$ and the molecular gas surface density $\\Sigma_\\mathrm{H2}$ for the $\\co$ complexes and $\\cotwo$ clouds.  The structures we have identified appear to follow the relation shown in \\citet{2009ApJ...699.1092H}, where $v_\\mathrm{o} \\propto \\Sigma_\\mathrm{H2}^{1/2}$ is expected for clouds in gravitational equilibrium, although these two variables are no independent ($\\Sigma_\\mathrm{H2} \\propto \\sigma_v$).    \\begin{figure*} \\centering \\begin{tabular}{ccc} \\epsfig{file= f14a.eps,width=0.3\\linewidth,angle=90} \\epsfig{file= f14b.eps,width=0.3\\linewidth,angle=90} \\epsfig{file= f14c.eps,width=0.3\\linewidth,angle=90} \\end{tabular} \\caption{Radial distributions of $\\Sigma_\\mathrm{H2}$ {\\bf (left)}, $\\Sigma_\\mathrm{SFR}$ {\\bf (center)} and $\\Sigma_\\mathrm{SFE}$ {\\bf (right)} of the regions observed over the disk, including both complexes and clouds.  Filled black dots represent on-arm complexes and clouds, while filled red dots illustrate the on-arm complexes and clouds located in regions 6, 7 and 11 (see Figures \\ref{figure_ngc6946e_13co} and \\ref{figure_map_sfe}).  Open circles illustrate inter-arm complexes and clouds.  Most of the on-arm structures that present higher $\\Sigma_\\mathrm{H2}$, $\\Sigma_\\mathrm{SFR}$ and $\\Sigma_\\mathrm{SFE}$ than the other complexes or clouds, particularly for the star formation rate, are located in regions 6, 7 and 11.  These structures are inside the region enclosed by $r=3.5$ kpc and $r=4.5$ kpc (illustrated by the black dashed lines, see Figure \\ref{figure_map_sfe}).} \\label{figure_rad_sfr} \\end{figure*}    How are the GMCs formed in non-grand design galaxies like NGC 6946? In flocculent galaxies (or galaxies with active potentials) the formation of inter-arm molecular gas structures is proposed to be the result of local instabilities rather than being the fragmented residuals left behind by the kinematic shear of the GMAs in spiral arms of grand design galaxies.  In this scheme, the properties of the clouds are approximately similar across the disk, and no offsets between gas and star formation are expected.  High resolution observations of flocculent galaxies have provided evidence in support of this scenario.  For instance, \\citet{2003PASJ...55..605T} mapped the $\\co$ over the southern arm of the flocculent galaxy NGC 5055.  They found no obvious offset between H$\\alpha$ and the molecular gas, and that on-arm and inter-arm clouds do not have significant differences in their properties.  Similarly, in the present study we have identified structures with similar properties across the observed portion of the disk.  Nevertheless, we have found that some of the on-arm clouds are more massive and have broader line widths than other clouds located in both on-arm and inter-arm regions.  Why are clouds departing from the mean properties observed elsewhere?  According to simulations of gas in galaxies with an active potential (\\citealt{2006MNRAS.371..530C}; \\citealt{2008MNRAS.385.1893D}; \\citealt{2011ApJ...735....1W}) the most massive structures are observed to be produced in regions where collision or merging between spiral arms occur, yielding an enhanced star formation in those overdense regions.  We have found that the more massive clouds spatially match sites of higher SFR, which may provides evidence in favor of these models.  A detailed description of the velocity field is needed to examine the potential presence of colliding flows of gas in this type of galaxy (see Section \\ref{velo-field} below).  \\begin{figure*} \\centering \\begin{tabular}{cc} \\epsfig{file=f15a.eps,width=0.4\\linewidth,angle=0} \\epsfig{file=f15b.eps,width=0.4\\linewidth,angle=0} \\end{tabular} \\caption{{\\bf Left}: Map of the Star Formation Efficiency derived for individual $\\co$ complexes.  The color bar is in units of Myr$^{-1}$.  The black contours highlight the complexes with $\\Sigma_\\mathrm{H2} > 110$ $\\Msun$ pc$^{-2}$.  {\\bf Right}: Map of the SFE for the $\\cotwo$ clouds.  In this case, black contours highlight clouds with $\\Sigma_\\mathrm{H2} > 135$ $\\Msun$ pc$^{-2}$.  As in the left panel, the color bar is in units of Myr$^{-1}$.  Circles illustrate regions where we found structures that deviate from the other identified structures in Figure \\ref{figure_rad_sfr}.  Dashed lines denote radii of $r=3.5$ kpc and $r=4.5$ kpc (see Figure \\ref{figure_rad_sfr}).} \\label{figure_map_sfe} \\end{figure*}   In Figure \\ref{figure_arm10} we found that some of the inter-arm $\\co$ complexes can have size and mass comparable to the most massive complexes found in on-arm structures.  However, taking advantage of the finer resolution provided by $\\cotwo$ observations, we have found that while the most massive on-arm clouds remain equally massive as we resolve the smaller structures, the inter-arm complexes are decomposed into several less massive components.  Nevertheless, the sizes, line widths, luminosities and masses are similar for most the structures observed in both on-arm and inter-arm regions.  Thus, the properties of the giant molecular clouds in NGC 6946 do not change substantially when they enter the spiral arms.  However, we observe two regions that show significant higher surface density and more massive GMCs than in other regions across the disk.  A natural comparison to our study of resolved GMCs in the disk of NGC 6946 may be made with the $\\co$ observations of the central 5 kpc of NGC 6946 presented in DM12. As shown in Figure \\ref{figure_line-width}a, despite differences in resolution and cloud identification algorithms, the sizes and velocity dispersions of the DM12 GMCs and our complexes (as traced by $\\co$) follow a similar trend (with the exception of the handful of DM12 clouds within 400 pc of the galaxy center), which is slightly steeper than that measured by S87 for Galactic disk GMCs. Figure \\ref{figure_line-width}b shows that the trends defined by the luminosities of the $\\co$ complexes and DM12 GMCs are also consistent. In Figure \\ref{figure_line-width}c, our observations indicate that a constant (i.e., not radially varying) value of $X_\\mathrm{CO}$ of $2 \\times 10^{20}$ is appropriate, which is consistent with the value typically assumed for the Milky Way disk clouds. This value is slightly higher than (but within the quoted errors of) DM12, who derive an average, non-radially varying value of $1.2 \\times 10^{20}$ in the central kiloparsecs of NGC 6946. We note that adopting an $I(2 \\rightarrow 1) = I(1 \\rightarrow 0)$ ratio of $< 1$ would reduce the best-fit value of $X_\\mathrm{CO}$ for the $\\cotwo$ clouds by the same factor.  Overall, the trend for the $\\co$ detections in Figure \\ref{figure_line-width}c is consistent between the two samples, but they begin to diverge at the highest cloud masses and luminosities, likely because the complexes presented here are blends of multiple GMCs.   \\subsection{Comparison with Galactic studies}\\label{discuss-gal} Recent studies of star formation rate (commonly traced by young stellar objects or YSOs) and column densities of gas (usually traced by near-IR extinction) in nearby Galactic clouds have compared Galactic and extragalactic gas and SFR surface density relations (\\citealt{2009ApJS..181..321E}; \\citealt{2010ApJ...723.1019H}; \\citealt{2011ApJ...739...84G}).  For instance,  \\citet{2010ApJ...723.1019H} found that the prescription given by  the Kennicutt-Schmidt law underpredicts the values found in nearby molecular clouds by factors that range from 21 to 54.  They found that such differences emerge from the different scales involved in calculating the SFR and the gas surface densities.  While the gas tracers used in Galactic studies are usually probing the denser component of the gas, the extragalactic observations have averaged over scales that include both dense and a more diffuse gas component which may not be related to the star formation process.  \\begin{figure} \\centering \\begin{tabular}{c} \\epsfig{file=f16.eps,width=0.8\\linewidth,angle=90} \\end{tabular} \\caption{Variation of the size-line width coefficient $v_\\mathrm{o}=\\sigma_v/R^{1/2}$ with the mass surface density of the structures derived in this study.  The dashed line represents the locus of clouds in virial equilibrium.} \\label{figure_co_fact} \\end{figure}   Figure \\ref{figure_sfr_heid} compares the values found in this work for NGC 6946 GMCs and the values presented by \\citet{2010ApJ...723.1019H} for Galactic star forming regions.  Besides the $\\Sigma_\\mathrm{gas}$ and $\\Sigma_\\mathrm{SFR}$ values for low-mass star formation regions traced by $A_V$ maps and counts of YSOs respectively, they included the most massive star-forming clumps that \\citet{2010ApJS..188..313W} studied with HCN gas maps.  We note that low-mass star-forming regions (red stars) overlap with the GMCs for $\\Sigma_\\mathrm{H2} \\sim 80\\ \\Msun\\ \\mathrm{pc}^{2}$.  Nevertheless, this apparent coincidence has to be considered carefully, as we are comparing low-mass (galactic) and massive star formation (extra-galactic).  Additionally, the masses of the large molecular clouds in \\citet{2010ApJ...723.1019H} study are two order of magnitudes smaller than the masses of the GMCs reported in this study.  Probably a more direct comparison to our GMCs would be provided by the massive HCN clumps, since extragalactic SFR tracers are exclusively sensitive to massive star formation.  In Figure \\ref{figure_sfr_heid} we have included the massive HCN clumps as red diamonds.  Surprisingly, if we extrapolate the $\\Sigma_\\mathrm{SFR}$-$\\Sigma_\\mathrm{H2}$ relation we found for $\\cotwo$ GMCs (Equation \\ref{sfr21_fit}) to higher gas surface densities, we observe that the massive dense clumps fall roughly along the relation.  Thus, the $\\Sigma_\\mathrm{SFR}$-$\\Sigma_\\mathrm{H2}$ relation for massive star formation regions is consistent with a quadratic relation.  Nevertheless, a more complete dynamical range in molecular gas and star formation surface density is needed to fill the gap between Galactic and extragalactic observations, and assess the intrinsic $\\Sigma_\\mathrm{SFR}$-$\\Sigma_\\mathrm{H2}$ relation.   \\subsection{Star formation and evolution of GMCs in NGC 6946}\\label{discuss-sfr} Although a complete unbiased survey of the GMCs population and their star formation activity in different environments is needed to properly investigate the $\\Sigma_\\mathrm{SFR}-\\Sigma_\\mathrm{H2}$ relation, our analysis allows us to shed light on the difference in star formation activity between on-arm and inter-arm clouds.  In Sections \\ref{sf-inon} and \\ref{rad-prop}, we observed a clear enhancement of SFR in on-arm clouds.  This enhancement is more pronounced for structures located in regions we suspect are the result of recent convergence of gas flows (regions 6, 7 and 11 in Figure \\ref{figure_ngc6946e_13co}).  In fact, it is in those regions where we find significant $\\cother$ emission, giving further observational evidence for the presence of denser gas in those structures.  We have found steeper slopes than previous studies on NGC 6946 for the $\\Sigma_\\mathrm{SFR}-\\Sigma_\\mathrm{H2}$ relation for both $\\co$ complexes and $\\cotwo$ clouds.  As was stated above, the slope of the relation $\\Sigma_\\mathrm{SFR}-\\Sigma_\\mathrm{H2}$ can be affected by several factors, including the method used to perform the linear regression between these two quantities, the tracer used to estimate gas densities and star formation activity, selection effects and the resolution of the maps used.  For the finest resolution in this paper, which is given by the $\\cotwo$ observations towards the brightest molecular gas regions, we have found a $\\Sigma_\\mathrm{SFR}-\\Sigma_\\mathrm{H2}$ relation that is almost quadratic.  A quadratic relation for $\\Sigma_\\mathrm{SFR}$-$\\Sigma_\\mathrm{H2}$ has been proposed by \\citet{2011ApJ...739...84G} for Galactic clouds.  They show that their data are consistent with a star formation law in which the $\\Sigma_\\mathrm{SFR}$ is proportional to the $\\Sigma_\\mathrm{gas}^2$.  Deviations from the power law are attributed to evolutionary stage differences in the local gas: in some older regions the gas can be disrupted by stars that have been formed there, and other younger regions can contain star clusters deeply embedded in dense gas at the onset of star formation.  \\begin{figure} \\centering \\begin{tabular}{c} \\epsfig{file=f17.eps,,width=0.8\\linewidth,angle=90} \\end{tabular} \\caption{Comparison between the $\\Sigma_\\mathrm{SFR}-\\Sigma_\\mathrm{H_2}$ values found in this work for NGC 6946 GMCs from $\\cotwo$ (black dots) and the values presented by \\citet{2010ApJ...723.1019H} for Galactic star forming regions (red symbols).  Stars show low-mass star formation regions, with $A_V$ maps and counts of YSOs used as tracers of gas and SFR respectively.  Diamonds illustrate the most massive star-forming clumps from \\citet{2010ApJS..188..313W} traced by HCN gas maps.  Red solid lines show the broken power law found by \\citet{2010ApJ...723.1019H}, while the black solid line illustrates the relation found in this present paper for $\\cotwo$ clouds.  Galactic low-mass star-forming regions overlap with the NGC 6946 GMCs for $\\Sigma_\\mathrm{H2} \\sim 80\\ \\Msun\\ \\mathrm{pc}^{2}$, although we emphasize that our estimates of SFR trace massive star formation.  Thus, HCN massive clumps represent a more direct comparison to our GMC values.  We observe that the massive dense clumps fall along a crude extrapolation of the $\\Sigma_\\mathrm{SFR}$-$\\Sigma_\\mathrm{H2}$ relation found for NGC 6946 GMCs (Equation \\ref{sfr21_fit}) to higher gas surface densities.} \\label{figure_sfr_heid} \\end{figure}   We noticed in Figure \\ref{figure_ngc6946e} that a few regions located in the outskirts of the disk present weak or undetectable emission in $\\co$ at our sensitivity level, but are bright in 24$\\mu$m, FUV, H$\\alpha$ and HI.  In order to detect the gas in some of those regions, we performed deeper low resolution observations of $\\co$ using CARMA in E array configuration towards the small region illustrated in Figure \\ref{figure_mom0_e}.  The sensitivity reached by these $\\co$ observations was $\\sigma \\sim$ 3 $\\Msun$ pc$^{-2}$ for a channel width of 2.5 $\\kms$ and resolution $8\\farcs45$ $\\times$ $7\\farcs29$ (a factor of $\\sim$ 4 higher than the sensitivity yielded by our extended mosaic).  We have applied the same procedure explained in Section \\ref{cprops} to identify discrete structures and estimate their properties.  Nevertheless, in this case we are only interested in the $\\Sigma_\\mathrm{H2}$, $\\Sigma_\\mathrm{HI}$ and the $\\Sigma_\\mathrm{SFR}$.  Figure \\ref{figure_sfr_e} shows the relation between these quantities for the complexes found in this region.  The right panel shows that these complexes exhibit low molecular gas content, but SFR comparable to denser molecular clouds observed in the complete mosaic (shown by open blue circles).  Additionally, in the left panel we observe that the atomic content is comparable to the molecular gas in these low CO-luminosity complexes.   A possible way to explain structures with high SFR and low molecular gas surface density is to consider an evolutionary scenario similar to that proposed by \\citet{2011ApJ...739...84G} for a sample of large Galactic molecular clouds.  Given that these molecular complexes in NGC 6946 appear to be isolated from the spiral arms, they could be GMCs that escaped from the gravitational potential minima, and consumed some of their molecular gas through forming stars.  That atomic gas peaks in these regions can be then naturally explained as a result of the photo-dissociation of the molecular gas by the newborn star population.  We observed peak atomic gas surface densities of $\\sim$ 35 $\\Msun$ pc$^{-2}$ in these regions, similar to the peaks found in the arm regions, and a factor of 3 larger than the azimuthally averaged value at that distance from the galactic center.  We postulate that, these structures deviate from the main $\\Sigma_\\mathrm{SFR}$-$\\Sigma_\\mathrm{H2}$ relation found for the remaining complexes or GMCs due to differences in evolutionary stage.  Higher sensitivity and resolution observations of the gas in these regions, as well as age estimates are needed to corroborate this scenario.  \\begin{figure} \\centering \\begin{tabular}{c} \\epsfig{file=f18.eps,width=0.8\\linewidth,angle=-90} \\end{tabular} \\caption{This figure illustrates the region of the eastern part of NGC 6946 we have observed with higher sensitivity using CARMA in E array.  The dashed line represent the 7-pointing mosaic we used to cover the region.  The contour map illustrates the $\\co$ intensity map, with contour levels of 2$^n$ $\\times$ 1.5 K $\\kms$ ($n \\ge 1$).  The color map shows the FUV+24$\\mu$m composite SFR map of the same region.  Red ellipses illustrate the molecular gas complexes with low $\\co$ emission, but associated SFR.} \\label{figure_mom0_e} \\end{figure}   \\subsection{Velocity field and comparison with models}\\label{velo-field} Although a detailed comparison between our observations and simulated non-grand design galaxies is beyond the scope of this paper, in this section we perform a simple comparison between the kinematics of the ISM observed in NGC 6946 and what is predicted by simulations.  In their simulation of a galactic disk with a live spiral model, \\citet{2011ApJ...735....1W} presented a dynamical picture of the behavior of the ISM.  They showed that both gas and the stellar arms roughly follow the local galactic rotational velocity, and they predicted that the relative velocities of the gas with respect to the mean rotational speed should be $\\lesssim$ 15 $\\kms$.  In this scenario, the cold gas does not follow the organized flow pattern predicted by the conventional density-wave theory and shows rather chaotic velocity structures around spiral arm potential minima, especially when the time dependence of stellar spiral potential is also calculated (\\citealt{2011ApJ...735....1W}). This effect was also observed by \\citet{2008MNRAS.385.1893D} in their simulation of spiral galaxies with an active potential.  In order to investigate the kinematics of gas flows in the ISM, we have created a velocity residual map in the region of NGC 6946 observed in this paper.  The velocity residual map was created by subtracting the local observed velocity field of the molecular gas traced by $\\co$ by a model of the global circular motion of the gas.  We aim to identify the non-circular motions of the molecular gas, and assess whether the places with significant velocity gradient are related to regions of dense gas and high star formation.  The model of the velocity field was created using the rotation curve derived by \\citet{2008AJ....136.2648D} for the HI map from THINGS.  We have used the GIPSY task {\\it velfi} to create the velocity field model from the parameters provided by \\citet{2008AJ....136.2648D} such as inclination, position angle, and the rotation velocity at a given radius using a tilted ring approach.  The final residual velocity map is shown in Figure \\ref{figure_mom1_e}.  We observe that the molecular gas shows irregular non-circular motions in the range $\\sim -10$ $\\kms$ to 15 $\\kms$ in the plane of the disk, in agreement with the simulation provided by \\citet{2011ApJ...735....1W}.  Moreover, we observe a region of strong velocity gradient coincident with one of the regions we found to present high SFE (region 11 in Figure \\ref{figure_map_sfe}).  In order to have a more detailed view of the velocity field in this region, in the left panel of Figure \\ref{figure_pv} we show a position-velocity diagram of region 11 taken parallel to the local velocity gradient.  The molecular gas shows a much steeper velocity gradient compared to the global rotation velocity described by the model.  According to models of spiral galaxies with an active potential, this region may be the result of the convergence of gas into the spiral potential from both sides, as we observe in the residual velocity field regions moving faster and slower than the global rotation model, i.e., a transition of positive to negative velocity values across the region.  \\begin{figure} \\centering \\begin{tabular}{ll} \\epsfig{file=f19a.eps,width=0.45\\linewidth,angle=90} & \\epsfig{file=f19b.eps,width=0.45\\linewidth,angle=90} \\end{tabular} \\caption{{\\bf Left}: Atomic gas surface density vs.\\ molecular gas surface density for $\\co$ complexes detected with higher sensitivity (red ellipses in Figure \\ref{figure_mom0_e}).  We observe that these complexes are located in regions where the gas is dominated by the atomic gas phase.  Symbols and labels are the same as Figure \\ref{figure_hi_h2}.  {\\bf Right}: Star formation rate vs.\\ molecular gas surface density for the $\\co$ complexes presented in the left panel (black circles).  As a comparison, blue open circles show the $\\co$ complexes found in the complete mosaic (left panel of Figure \\ref{figure_sfr}).  The complexes observed with higher sensitivity appear to deviate from the trend found by \\citet{2008AJ....136.2846B} due to their high star formation activity but low molecular gas content, as is shown in Figure \\ref{figure_mom0_e}.  These cases may represent complexes that have been consuming their molecular gas by forming stars, yielding HI from photodissociation of H$_2$ by the ambient radiation field.} \\label{figure_sfr_e} \\end{figure}   The other region with high SFE (region 6 in Figure \\ref{figure_ngc6946e_13co}), on the other hand, does not present such a high velocity gradient as the region discussed above.  However, we observe a $\\sim$15 $\\kms$ difference in velocity between region 6 and the regions located to the north-west and to the south-east.  While those regions surrounding the molecular gas complex in region 6 are moving more or less at $-5$ $\\kms$ with respect to the rotating frame, region 6 is moving in the other direction with a velocity of $\\sim$10 $\\kms$.  This difference in velocity can be observed clearly in the position-velocity diagram shown in the right panel of Figure \\ref{figure_pv}.  This may imply that the gas actively forming stars in region 6 could be the result of a recent convergence of flows.  Although we do not observe signatures of a merging or collision process of spiral arms in the gas distribution near region 6, it is possible that the process had already formed a new spiral arm, obscuring the original arm locations.  More detailed analysis of the kinematics of the gas is required to disentangle the actual process involved in the enhancement of star formation efficiency in region 6.  \\begin{figure} \\centering \\begin{tabular}{c} \\epsfig{file=f20.eps,width=0.85\\linewidth,angle=0} \\end{tabular} \\caption{Residual velocity field of the region observed in NGC 6946.  The color bar is in units of $\\kms$.  Black solid lines illustrate the slices used to generate the position-velocity diagrams shown in Figure \\ref{figure_pv}.  We have subtracted a model of the rotation velocity from the observed $\\co$ velocity field.  The model of the rotation velocity was created using the rotation curve derived by \\citet{2008AJ....136.2648D}.   We observe that there is a significant velocity gradient in one of the northern regions where we detect high SFE (region 11 in Figure \\ref{figure_map_sfe}).  In the southern region (region 6 and 7 in Figure \\ref{figure_map_sfe}), on the other hand, we do not observe a significant velocity gradient inside the individual CO complexes, but there is a $\\sim$ 15 $\\kms$ discontinuity between the gas in regions 7 and 6.} \\label{figure_mom1_e} \\end{figure}   \\subsection{Caveats of this work}\\label{caveats} The study presented here can be improved in multiple ways.  Our analysis has been limited to the north-eastern area of the molecular gas disk of NGC 6946 for $\\co$ observations, and to some of the brightest regions for $\\cotwo$ follow-up observations.  A full census of the GMCs population across the disk is needed in order to perform a detailed study of molecular cloud properties and their relation with the surrounding environment.  Nevertheless, high sensitivity maps with resolutions $\\sim 60$ pc remain observationally challenging to obtain, although recent studies have achieved this goal with CARMA+NRO 45m (e.g. DM12).  New facilities like ALMA will provide maps of unprecedented quality, allowing us to perform an unbiased analysis of the molecular gas in nearby galaxies at scales close to GMC sizes.  In our SFR calculations for $\\cotwo$ clouds, we have applied a crude global extinction correction A$_\\mathrm{H\\alpha}=1$ mag to the SFR(H$\\alpha$).  Although we showed that such correction is a good approximation to the SFR values found using a combination of FUV+24$\\mu$m for $\\co$ complexes, a more sophisticated approach to calculate the obscured SFR at resolution close to the GMCs sizes is needed.  Thus, tracers relatively unaffected by dust absorption represent a valuable alternative to estimate the total SFR at smaller scales.  In this direction, observations of Paschen $\\alpha$ (P$\\alpha$) line could provide unbiased measurements of the number of ionizing photons in these star forming regions (\\citealt{2007ApJ...666..870C}).  Equations \\ref{sfr} and \\ref{sfr_ha} assume that stars have been formed continuously.  Nevertheless, this assumption may not be longer valid for scales close to GMC sizes, as one approaches the case of  a single stellar population with a discrete age.  Some studies have investigated the uncertainty in using a linear conversion from UV or H$\\alpha$ luminosity to estimate the SFR, making use of models assuming an instantaneous burst of star formation (\\citealt{2010MNRAS.407.2091G}; \\citealt{2012AJ....144....3L}).  \\citet{2012AJ....144....3L} showed that while H$\\alpha$ emission occurs mostly in the first 10 Myr, FUV presents low, but significant, emission up to 65 Myr after the burst.  They report a factor of 2 uncertainty inferring the SFR from H$\\alpha$, and a factor 3-4 uncertainty inferring SFR from FUV.  Resolved star cluster observations in the regions observed in this study could shed light on the ages involved, and provide an independent way to estimate the star formation rate which can be compared to the SFR traced by FUV or H$\\alpha$.  \\begin{figure} \\centering \\begin{tabular}{cc} \\epsfig{file=f21a.eps,width=0.45\\linewidth,angle=0} \\epsfig{file=f21b.eps,width=0.5\\linewidth,angle=0} \\end{tabular} \\caption{Position-velocity diagrams of the regions showing velocity features in the residual velocity field.  The extent and orientation of the slices are shown by the black solid lines in Figure \\ref{figure_mom1_e}.  Color map shows the $\\co$ data, and the black contours illustrate the model of the rotation velocity created using the rotation curve by \\citet{2008AJ....136.2648D}.  {\\bf Left}:  Position-velocity diagram of region 11.  We observe that the velocity gradient is sharper in the $\\co$ data than the global rotation velocity model.  {\\bf Right}:  Position-velocity diagrams of the regions 6 (peak on right) and 7.  We do not observe a significant velocity gradient inside the individual CO complexes, but the molecular gas in region 6 is moving about $\\sim$ 15 $\\kms$ away from the contiguous region 7.} \\label{figure_pv} \\end{figure}"
    },
    "1208/1208.1273_arXiv.txt": {
        "abstract": "We present the discovery and phase-coherent timing of four highly dispersed millisecond pulsars (MSPs) from the Arecibo PALFA Galactic plane survey: PSRs J1844+0115, J1850+0124, J1900+0308, and J1944+2236. Three of the four pulsars are in binary systems with low-mass companions, which are most likely white dwarfs, and which have orbital periods on the order of days. The fourth pulsar is isolated. All four pulsars have large dispersion measures (DM $> 100$ pc cm$^{-3}$), are distant ($\\ga 3.4$ kpc), faint at 1.4 GHz ($\\la 0.2$ mJy), and are fully recycled (with spin periods $P$ between 3.5 and 4.9 ms). The three binaries also have very small orbital eccentricities, as expected for tidally circularized, fully recycled systems with low-mass companions.  These four pulsars have DM/$P$ ratios that are among the highest values for field MSPs in the Galaxy.  These discoveries bring the total number of confirmed MSPs from the PALFA survey to fifteen.  The discovery of these MSPs illustrates the power of PALFA for finding weak, distant MSPs at low-Galactic latitudes. This is important for accurate estimates of the Galactic MSP population and for the number of MSPs that the Square Kilometer Array can be expected to detect. ",
        "introduction": "The PALFA survey is an ongoing, large-scale pulsar survey of the Galactic plane that uses the Arecibo 305-m radio telescope and the Arecibo L-Band Feed Array (ALFA) 7-beam multi-beam receiver \\citep{cfl+06}.  It is one of Arecibo's key science projects, and it will ultimately cover the entire Arecibo-visible sky within $5^{\\circ}$ of the Galactic plane (longitudes of $32^{\\circ} \\la l \\la 77^{\\circ}$ and $168^{\\circ} \\la l \\la 214^{\\circ}$).  PALFA observes at relatively high observing frequencies ($1220-1520$\\,MHz) in order to mitigate the deleterious effects that interstellar dispersion and scattering have on the detection of distant pulsars at low Galactic latitudes.  In this sense, the PALFA survey is similar to the highly successful Parkes Multibeam Pulsar Survey (PMPS; Manchester et al. 2001\\nocite{mlc+01}), but with increased time and frequency resolution, such as those of the current Parkes HTRU survey \\citep{kjv+10}.  Compared to past Arecibo surveys, PALFA explores a far larger spatial volume due to its high time and frequency resolution, enabling discovery of faint, highly dispersed millisecond pulsars (MSPs)\\footnote{An MSP typically refers to a (partially or fully) recycled pulsar having a small surface magnetic field strength ($B \\la 10^{10}$ G) and a large characteristic age ($\\ga 10^{9}$ yr).}  in the Galactic plane, which tend to be at larger distances.  Since the majority of MSPs have binary companions ($\\sim 80$\\%, Lorimer 2008\\nocite{l08}), they are often interesting test cases for studies of exotic binary stellar evolution (e.g., Archibald et al. 2009\\nocite{asr+09}; Freire et al. 2011\\nocite{fbw+11}).  The unsurpassed sensitivity of the Arecibo telescope is highly advantageous for detecting binary MSPs because the PALFA pointing dwell times are only a small fraction (less than 10\\%) of the orbital periods of all known binary radio pulsars.  In this regime, linear acceleration searches are highly effective at recovering Doppler-smeared periodicities \\citep{jk91}.  The discovery of PSR J1903+0327 by \\citet{crl+08} is an excellent illustration of the PALFA survey's sensitivity to highly dispersed pulsars.  PSR J1903+0327 has a spin period of only 2.15 ms and has the highest dispersion measure (DM) of all known completely recycled Galactic MSPs, 297 pc cm$^{-3}$.  It occupies a region of DM-period phase-space that has previously been unexplored for pulsars in the Galactic field.  For example, of the 90 recycled Galactic field radio pulsars currently listed in the ATNF catalog (apart from the PALFA discoveries)\\footnote{We also exclude the high-DM MSPs in globular clusters, because these are found by targeted searches conducted with very long integration times and at even higher observing frequencies (generally $\\sim 2$ GHz).}  with spin periods $P < 25$ ms, only 9 have DM $> 100$ pc cm$^{-3}$, and only 14 have DM-derived distances listed in the catalog that are greater than 3 kpc.  One of the main motivations for finding distant, highly dispersed MSPs is to determine a more complete census of the Galactic MSP population, which is currently biased by the large number of nearby ($\\la 2$ kpc) sources -- especially given the recent discoveries of more than 40 generally nearby MSPs through targeted searches of {\\it Fermi} gamma-ray sources (e.g., Ransom et al. 2011\\nocite{rrc+11}).  MSPs are the longest-living active manifestations of neutron stars, with active lives hundreds to thousands of times longer than those of normal pulsars, magnetars or accreting neutron stars.  Hence they give valuable insight into the Galactic neutron star population and binary stellar evolution in particular.  The planned Square Kilometer Array (SKA; e.g., Carilli \\& Rawlings 2004\\nocite{cr04}) should be able to detect a large fraction of the MSPs in the Galaxy, and finding high-DM MSPs beforehand will help tell us how many to expect.  Furthermore, MSPs are excellent probes of the interstellar medium, and the discovery of more distant MSPs whose signals likely pass through multiple scattering screens opens new possibilities (and challenges) in this area. For instance, scattering measurements of large-DM MSPs can be used to compare the observed effects of scattering on timing behavior with predictions of those effects.  MSPs can also be used as high-precision astronomical clocks.  Stable MSPs form the basis of the pulsar timing efforts of the International Pulsar Timing Array consortium to detect long-period gravitational radiation from observations of pulsar timing residuals \\citep{jhl+05,h++10}.  The recent precision mass measurements of PSR J1903+0327 \\citep{fbw+11} and PSR J1614$-$2230 \\citep{dpr+10} indicate that MSPs can have masses well above the Chandrasekhar mass, and mass constraints from new MSPs will continue to map out the MSP mass distribution. The PALFA pulsar PSR J1949+3106, which may also have a higher mass than the Chandrasekhar mass \\citep{dfc+12}, is an example where future mass measurements may prove important.  Although high-DM MSPs have timing precision problems associated with interstellar scattering in addition to their being generally radio-faint, future observations with next generation instruments like the Square Kilometer Array (SKA) may be able to mitigate these factors. Precision measurement of NS masses (e.g., Demorest et al. 2010\\nocite{dpr+10}) and the measurement of ultra-high spin rates (e.g., Lattimer \\& Prakash 2007; Hessels et al. 2006\\nocite{lp07, hrs+06}) can also rule out some high-density NS equations of state.  In this paper we report the discovery and follow-up timing of four highly dispersed MSPs from the PALFA survey: PSRs J1844+0115, J1850+0124, J1900+0308, and J1944+2236.  These MSPs all have DMs that are in the top 5\\% of non-PALFA field radio MSPs.  Three of the four pulsars are in binary systems with low-mass companions.  These discoveries bring the number of PALFA recycled pulsar discoveries to 15, including PSR J1903+0327 \\citep{crl+08}, PSRs J1949+3106 and J1955+2427 \\citep{dfc+12}, and the partially recycled PSRs J2007+2722 and J1952+2630 \\citep{kac+10,kla+11}. The six additional recycled pulsars that have been confirmed in the survey need to be timed further to establish their rotational and orbital characteristics. These six pulsars will be published in forthcoming papers and are not considered further here.  In \\S2 we describe the discovery and follow-up observations of the four pulsars, and we present their phase-coherent timing solutions.  In \\S3 we discuss our results, and in \\S4 we present our conclusions. ",
        "conclusions": "With its high time and frequency resolution and relatively high observing frequency, the PALFA survey is sensitive to Galactic field MSPs at large distances and DMs.  The survey has so far discovered a total of 15 MSPs with $P < 25$ ms, 14 of which have DM $> 100$ pc cm$^{-3}$. Four of these PALFA MSPs (PSRs J1844+0115, J1850+0124, J1900+0308, and J1944+2236) are presented here with phase-coherent timing solutions. As an ensemble, the PALFA MSP discoveries show the ability of PALFA to extend the volume of MSP discovery space to relatively high DMs.  This is the first step toward a nearly complete census of the Galactic MSP population, which will be possible with the SKA."
    },
    "1208/1208.1790_arXiv.txt": {
        "abstract": "The Fermi observations of GRB 110721A \\citep{Fermi+12epeak} have revealed an unusually high peak energy $\\sim 15$ MeV in the first time bin of the prompt emission. We find that an interpretation is unlikely in terms of internal shock models, and confirm that a standard black-body photospheric model also falls short. We show that dissipative photospheric synchrotron models  ranging from extreme magnetically dominated to baryon dominated dynamics, on the other hand, are able to accommodate such high peak values. ",
        "introduction": "\\label{sec:intro} High energy ($>$ 100 MeV) observations of gamma-ray bursts (GRBs) have gained renewed interest since the launch of Fermi LAT \\citep{Atwood+09LAT}.  This instrument has revealed a curious diversity in the high-energy behavior of GRBs and consequently represents a new challenge for models  \\citep[see ][ for a recent review]{Mesz12GRBCTA}.  For GRB 110726A \\citep{Fermi+12epeak},\tthe extremely high $\\ve_{\\rm peak}\\approx 15 \\MeV$ peak energy (peak of the $\\ve F_\\ve$ spectrum),   which was observed right after the onset of the burst, makes its interpretation in the framework of a simple internal shock synchrotron model challenging. Interpreting the peak as a modified black-body from a simple photosphere is also difficult, as it can only account for energies up to few MeVs, and even the subsequent time-bins show  peak fluxes which, even though lower, are still unusually high. Such black-body photosphere models\thave been considered by \\citet{Fan+12corr} to explain a $L\\propto\\ve_{\\rm peak}^{0.4}$ correlation found  by \\citet{Ghirlanda+12comoving},  but as shown by \\citet{B-Zhang+12}, they must lie below a line in the $\\ve_{\\rm peak} -L$ plane which excludes the observed values for GRB 110721A.  Here we show that the high peak energy and luminosity of GRB110721A can be interpreted in the framework of a different class of photospheres, where dissipation occurs near the photospheric radius, and the spectrum is characterized by a synchrotron peak. Dissipative synchrotron photospheres have been considered for baryonic dominated regimes by, e.g.  \\citet{Rees+05photdis, Peer+06phot, Beloborodov10phot} and for magnetically dominated regimes by \\citet{Giannios06flare, Giannios12peak}.\tHere we consider such photospheres for an arbitrary acceleration law, characteristic of either a baryonic or a magnetic dominated regime.  Also the observed thermal bump close to $100 \\kev$ is naturally incorporated in this model. ",
        "conclusions": "\\label{sec:disc} We have considered the high peak energy values of the prompt spectrum of GRB 110721A, which reach as high as $\\ve_{\\rm peak}\\sim 15$ MeV \\citep{Fermi+12epeak}. A consideration of the usual internal shock prompt emission spectrum model shows that such high values are unlikely in this model. Furthermore, we confirm the conclusion of \\citet{B-Zhang+12} that a (non-dissipative) standard black-body baryonic photosphere model also cannot explain such high peak values and fluxes. However, we show that dissipative photospheric models, with a typical peak energy due to synchrotron radiation, are able to accommodate such high peak energies and flux values, with reasonable parameters, for cases where the dynamics is either baryonically or magnetically dominated.  If the temperature of the putatively observed black-body component can be used as a discriminant, this would seem to favor a more baryon dominated dissipative photosphere model, $\\mu$ closer to $1$, although a moderate magnetically dominated photosphere may also be possible.  PV and PM acknowledge NASA NNX09AL40G and OTKA grant K077795 for partial support, BBZ acknowledges the support from NASA SAO SV4-74018.  We thank Bing Zhang, Rui-Jing Lu, En-Wei Liang and Xue-Feng Wu for supplying the time binned data and for discussing their results { and the referee for a thorough report. }"
    },
    "1208/1208.3795_arXiv.txt": {
        "abstract": "Of the over 800 exoplanets detected to date, over half are on non-circular orbits, with eccentricities as high as 0.93.  Such orbits lead to time-variable stellar heating, which has major implications for the planet's atmospheric dynamical regime. However, little is known about the fundamental dynamical regime of such planetary atmospheres, and how it may influence the observations of these planets. Therefore, we present a systematic study of hot Jupiters on highly eccentric orbits using the SPARC/MITgcm, a model which couples a three-dimensional general circulation model (the MITgcm) with a plane-parallel, two-stream, non-grey radiative transfer model.  In our study, we vary the eccentricity and orbit-average stellar flux over a wide range.  We demonstrate that the eccentric hot Jupiter regime is qualitatively similar to that of planets on circular orbits; the planets possess a superrotating equatorial jet and exhibit large day-night temperature variations.  As in Showman and Polvani (2011), we show that the day-night heating variations induce momentum fluxes equatorward to maintain the superrotating jet througout its orbit.  We find that as the eccentricity and/or stellar flux is increased (corresponding to shorter orbital periods), the superrotating jet strengthens and narrows, due to a smaller Rossby deformation radius.  For the cases that are most distant and slowly rotating, we find hints of a regime shift, with no eastward flow at lower pressures.  For a select number of model integrations, we generate full-orbit lightcurves and find that the timing of transit and secondary eclipse viewed from Earth with respect to periapse and apoapse can greatly affect what we see in infrared (IR) lightcurves; the peak in IR flux can lead or lag secondary eclipse depending on the geometry.  For those planets that have large temperature differences from dayside to nightside and rapid rotation rates, we find that the lightcurves can exhibit ``ringing\" as the planet's hottest region rotates in and out of view from Earth.  These results can be used to explain future observations of eccentric transiting exoplanets. ",
        "introduction": "Since the first planetary confirmations in the mid-1990s \\citep{mayor1995,wolszczan1994} the detection and characterization of extrasolar planets continue to be major fields in astronomy and planetary science.  Over 700 planets have been detected from the ground and space, more than half of which are classified as ``hot Jupiters\", Jovian-mass planets that orbit their parent stars at distances less than 0.1 AU.  A number of these hot Jupiters transit their host star along our line of sight, allowing us to observe them as they pass in front and behind their parent star \\citep[e.g.][]{charb2008,swain2008,pont2008,knutson2007,knutson2008,knutson2009}. Using these observations, we can infer much about these planets' atmospheric composition, temperature structure, and circulation.  A fifth of these transiting exoplanets have eccentricities greater than 0.1, with values as large as 0.93 \\citep[HD80606b,][]{naef2001}.  These eccentric planets are subject to highly time-variable heating which has a significant effect on the planet's atmospheric dynamics.  Among planets amenable to observational follow-up, HAT-P-2b, which has an eccentricity of 0.52, undergoes a factor of 9 variation in flux throughout its orbit.  HD 17156b ($e=0.67$) experiences a factor of 27 variation in stellar flux, and HD 80606b, with its large eccentricity ($e=0.93$), undergoes an impressive factor of 828 variation in flux.  Because they are transiting, we can probe their atmospheres as we can for planets on circular orbits.  However, while it is clear the strongly variable heating leads to a vastly different {\\it forcing} regime than for exoplanets on circular orbits, it remains unknown whether this causes a fundamentally different {\\it dynamical} regime: is the circulation quantitatively similar to that on circular hot Jupiters, or is it a completely new circulation regime?  Eccentric transiting exoplanets present unique challenges when one attempts to extract information about their atmospheres from observational data.  In particular, interpretation of flux maxima and minima in infrared lightcurves can be complicated by the convolution of spatial effects (for example, hot spots on the planet that rotate into and out of view along our line of sight) with temporal effects (planet getting colder/warmer at apoapse/periapse passage).  Langton and Laughlin (2008a,b) and Cowan and Agol (2011) address this problem as applied to particular targets, but use only a two-dimensional hydrodynamical model and one-dimensional semi-analytic model, respectively.  To fully capture these effects, a three-dimensional circulation model that self-consistently calculates heating and wind velocities is needed.  Hence, it is crucial to conduct a comprehensive study that establishes the dynamical regime, temperature structure, and observational implications of eccentric exoplanets.  We use a three-dimensional atmospheric circulation model coupled to a non-grey radiative transfer scheme to study eccentric hot Jupiters as a whole.  In Section 2 we will describe our model setup and integrations.  Section 3 describes the dynamical regime and its dependence on eccentricity and mean stellar flux.  Section 4 presents synthetic lightcurves and attempts to determine what can be learned about atmospheric circulations from remote measurements.  Finally, Section 5 concludes and compares results to known eccentric hot Jupiters. ",
        "conclusions": "We present three-dimensional circulation models coupled with a two-stream, non-grey radiative transfer scheme for a number of theoretical eccentric hot Jupiters.  We have shown that as in published models with zero eccentricity, our high-eccentricity circulation models are dominated by eastward flow at photospheric levels which cause an eastward displacement of the hottest regions from the substellar point.  The rapid rotation rates associated with pseudo-synchronization at high eccentricity lead to a small Rossby deformation radius and in some cases multiple jets in the atmosphere.  Global-mean temperatures and day-night temperature differences peak not at periapse but several hours afterward due to finite radiative timescales in the planet's atmosphere.  Furthermore, we show that equatorial superrotation is generated and maintained by eddies formed by the strong day-night heating contrast, which induce a flux of momentum from midlatitudes to the equator.  The eddy magnitudes and momentum fluxes peak just after periapse passage leading to variations in the zonal-mean flow throughout the orbit.  Lastly, we have shown that the spatial and temporal variations of the wind and temperature structure, as well as the orbital viewing geometry of the system with respect to Earth, can affect the time and amplitude of peak IR flux seen in full-orbit lightcurves.  Depending on the viewing geometry of the orbit relative to Earth, we find that peaks in IR flux that either lead or lag periapse are possible; in all cases, a combination of temporal effects (temperatures changing over time) and geometric effects (hot spots rotating into or out of view) are important in controlling the timing and amplitude of the flux peaks.  In cases where the day-night temperature contrast is large and the rotational period is short, the lightcurve can also exhibit ``ringing\" in flux as the hottest region of the planet rotates in and out of view.  This ringing is non-periodic, due to the variation in stellar heating as a function of distance."
    },
    "1208/1208.3937_arXiv.txt": {
        "abstract": "We point out a possible generation mechanism of non-Gaussian bubbles in the sky due to bubble nucleation in the early universe. We consider a curvaton scenario for inflation and assume that the curvaton field $\\phi$, whose energy density is subdominant during inflation but which is responsible for the curvature perturbation of the universe, is coupled to another field $\\sigma$ which undergoes false vacuum decay through quantum tunneling. For this model, we compute the skewness of the curvaton fluctuations due to its interaction with $\\sigma$ during tunneling, that is, on the background of an instanton solution that describes false vacuum decay. We find that the resulting skewness of the curvaton can become large in the spacetime region inside the bubble. We then compute the corresponding skewness in the statistical distribution of the cosmic microwave background (CMB) temperature fluctuations. We find a non-vanishing skewness in a bubble-shaped region in the sky. It can be large enough to be detected in the near future, and if detected it will bring us invaluable information about the physics in the early universe. ",
        "introduction": "\\label{sec:introduction} Inflation, a stage of accelerated expansion in the very early universe, is now widely accepted as part of the standard evolutionary scenario of the universe. On the other hand, many models of inflation have been proposed but we are still far from being able to narrow down the possible models sufficiently. Among those many models of inflation, much attention has been paid recently to the ones based on string theory~\\cite{Kachru:2003aw}, which is considered to be a promising candidate for the ultimate unified theory. In particular, it is of great interest if the string theory landscape~\\cite{Susskind:2003kw}, in which there are many local minima, or false vacua, and the universe jumps from one minimum to another by quantum tunneling, can be observationally tested~\\cite{Yamauchi:2011qq,Yamamoto:1996qq}. Quantum tunneling of a scalar field with gravity is usually treated with the Coleman-De Luccia (CDL) instanton method~\\cite{Coleman:1980aw}, where the evolution of a scalar field is described with an O(4)-symmetric instanton, which is a solution of the Euclidean equations of motion. Motivated by the string landscape,  inflation models with  multi-scalar fields and/or with tunneling are now keenly studied. Extension of the CDL instanton method to a multi-scalar field system is already discussed in \\cite{Sugimura:2011tk}. In those studies, however, many inflation models have been proposed, and now it is important to distinguish those models by observation.  Observations of the power spectrum of CMB temperature anisotropies have convinced us of the existence of an inflationary stage in the very early universe. On top of that, if a non-Gaussian feature such as skewness or bispectrum is detected in the CMB anisotropy, it will have a strong impact on the physics of the early universe~\\cite{Komatsu:2010fb}. In particular, since any single-field slow-roll inflation model with canonical kinetic term predicts almost Gaussian fluctuations~\\cite{Maldacena:2002vr}, any non-zero non-Gaussianity will exclude all these models. Observations of non-Gaussianity use templates, such as local type~\\cite{Komatsu:2001rj},\\footnote{ In the view point of statistical distribution, local type non-Gaussianity, which is local in the sense of generation mechanism, is homogeneous and isotropic.} equilateral type~\\cite{Creminelli:2005hu}, and orthogonal type~\\cite{Senatore:2009gt}, in order to increase the statistical significance. However, since all of these templates assume statistical isotropy, there may be anisotropic non-Gaussianities which may not have been detected by these templates.  In this letter, we study a multi-field model in which a nonlinear interaction between two scalar fields, one of which being responsible for the curvature perturbation of the universe (that is, for the formation of the large scale structure) and the other for quantum tunneling via a CDL instanton, induces an anisotropic non-Gaussianity. To be specific, we introduce an inflaton field $\\Phi$ that realizes slow-roll inflation, a tunneling field $\\sigma$ that governs the tunneling dynamics, and a curvaton field $\\phi$ that contributes to the curvature perturbation of the universe~\\cite{Lyth:2002my,Sasaki:2006kq}.  The inflaton dominates the energy density of the universe during inflation but rapidly decays to radiation after the end of inflation, On the other hand, the energy density of the curvaton field is negligible during inflation but its decay is delayed after inflation so that it gradually begins to dominate the universe.  Approximating the universe during inflation by an exact de Sitter spacetime, the inflaton behaves as a cosmic clock and determines an appropriate time-slicing, namely, a spatially flat time-slicing of the de Sitter spacetime. In this setup, assuming that the energy scale associated with the tunneling field is much smaller than the energy scale of inflation, the bubble nucleation can be well described by a single-field CDL instanton with no backreaction to the geometry, that is, on the exact de Sitter background~\\cite{Coleman:1980aw,Yamamoto:1996qq}.  The curvaton $\\phi$ is affected by the background bubble through a coupling with the tunneling field $\\sigma$. For simplicity and definiteness, we consider a potential of the form, $V_\\mathrm{int}(\\sigma ,\\phi )=\\tilde\\lambda (\\sigma )\\phi^3$\\,. Assuming that $V_\\mathrm{int}(\\sigma, \\phi )$ is non-vanishing only at or inside the bubble wall, we expect that $\\phi$ may have a spatially localized, bubble-shaped non-Gaussianity due to the background bubble-shaped configuration of $\\sigma$. This leads to an anisotropic, bubble-shaped skewness of the CMB temperature anisotropy.  This paper is organized as follows. We first illustrate the background spacetime and the configuration of the bubble. Next, we briefly review a useful formalism for computing the equal-time $N$-point functions, the tunneling in-in formalism, and calculate the skewness of the curvaton fluctuations. Then we demonstrate that a sky map of an anisotropic non-Gaussian parameter $f_{\\rm NL}$ in our model. Finally, we end with a few concluding remarks. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we calculated the skewness in the CMB temperature anisotropy in a model with bubble nucleation during inflation, motivated by the string theory landscape. We considered bubble nucleation in the curvaton scenario of inflation in which the curvaton vacuum fluctuations are affected by the bubble nucleation through interaction with the tunneling field. The calculation was done by extending the in-in formalism to the instanton background~\\cite{preparation}. We found that there can be spatially localized, bubble-shaped skewness which is large inside the bubble.  As far as we know, bubble-shaped non-Gaussianities have not been studied carefully yet in observation. So it seems interesting to look for such a non-Gaussianity already in the current observational data. As pointed out by Komatsu and Spergel~\\cite{Komatsu:2001rj}, the bispectrum, corresponding to the 3-point function, contains much more information than the single value of the skewness. Thus, analysis beyond skewness may improve the observability of bubble-shaped non-Gaussianities. We hope to come back to this issue in a future publication.  Since the string theory landscape gives a strong motivation for inflation models with bubble nucleation, studies of such inflation models may be regarded as testing string theory using the universe as a laboratory. In any case, if any signature of bubble nucleation during inflation is found in observation, it gives a huge impact on the physics of the early universe, including string theory.  Finally, we note that in the evaluation of the skewness we neglected the effect of deviations from the Bunch-Davies vacuum. This may affect details of our result, though generic features are expected to remain the same. This effect can be evaluated by studying the evolution of mode function on the instanton background~\\cite{Yamamoto:1996qq}. We plan to come back to this issue in the near future."
    },
    "1208/1208.6515_arXiv.txt": {
        "abstract": "{We  analyze Chandra  observations  of  diffuse soft  X-ray emission associated with a complete  sample of 3CR radio galaxies at z $<$ 0.3. In this paper we focus on the properties of the spectroscopic sub-classes  of  high  excitation  galaxies (HEGs)  and  broad  line objects (BLOs).  Among the 33 HEGs we detect extended (or possibly extended) emission  in about  40\\% of the  sources; the  fraction is even  higher (8/10)  restricting the  analysis to  the  objects with exposure times larger than 10  ks. In the 18 BLOs, extended emission is seen only in 2 objects; this lower detection rate can be ascribed to the  presence of their  bright X-ray nuclei that  easily outshine any genuine diffuse emission.   A very close correspondence between  the soft X-ray and optical line morphology emerges. We also find  that the ratio between [O~III] and extended soft X-ray  luminosity  is confined within  a factor  of 2 around a median value of 5.  Both results are similar to what is seen in Seyfert galaxies.  We  discuss different  processes  that could  explain  the soft  X-ray emission  and  conclude  that  the photoionization  of  extended  gas, coincident with the narrow line region, is the favored mechanism.} ",
        "introduction": "\\begin{figure*} \\centering \\includegraphics[width=4.5cm]{19561f1.eps} \\includegraphics[width=4.5cm]{19561f2.eps} \\includegraphics[width=4.5cm]{19561f3.eps} \\includegraphics[width=4.5cm]{19561f4.eps} \\includegraphics[width=4.5cm]{19561f5.eps} \\includegraphics[width=4.5cm]{19561f6.eps} \\includegraphics[width=4.5cm]{19561f7.eps} \\includegraphics[width=4.5cm]{19561f8.eps} \\includegraphics[width=4.5cm]{19561f9.eps} \\includegraphics[width=4.5cm]{19561f10.eps} \\includegraphics[width=4.5cm]{19561f11.eps} \\includegraphics[width=4.5cm]{19561f12.eps} \\includegraphics[width=4.5cm]{19561f13.eps} \\includegraphics[width=4.5cm]{19561f14.eps} \\includegraphics[width=4.5cm]{19561f15.eps} \\includegraphics[width=4.5cm]{19561f16.eps} \\caption{Images in the  soft X-ray band for the  16 sources classified as extended  or possibly extended  in Tab. \\ref{bigtable}  (plus the complex source 3C~321 and the \"blobby\" source 3C~277.3).  The images are unbinned  (the pixel  size is 0.49\\arcsec)  and we  indicate the orientation of  the radio axis.  When extended  emission is present, we superposed the region used for the spectral extraction.} \\label{images} \\end{figure*}  Supermassive  black  holes  (SMBHs)  have  a profound  effect  on  the evolution of galaxies but the nature of the relationship between these two entities is still an open problem. This depends on how much of the released  energy interacts with  the surrounding  matter and  how this accretes onto the SMBHs. This feedback process can be explored through a  multi-wavelength   analysis  of   the  emission  observed   in  the circumnuclear regions. The vast  collection of emission lines, ranging from  the mm  to the  X-ray bands,  reveal the  presence of  a complex multiphase  medium surrounding the  AGN. Of  special interest  for our purposes  is the  so-called  Narrow Line  Region  (NLR) where  optical emission  lines with  widths of  several hundreds  of km  s$^{-1}$ are produced.  The NLR  is just outside (or even  within) the SMBHs sphere of influence  and its physical  and dynamical properties  are strongly affected  by the  central  engine.  It  exists  at the  interface between the  active nucleus and the  galaxy, and is  thus a convenient laboratory  in  which  to  explore  the energy  exchange  between  the two.  Moreover,  the NLR  is  mostly  free  from the  effects  of obscuration  and it is  resolved in  most nearby  AGN, allowing  us to perform  spatially resolved studies  and morphological  comparisons in different energy bands. Indeed, it offers a wide variety of diagnostic tools to probe the gas physical conditions.  The  NLR and  the processes  that occur  in it  have  been extensively studied  in many  bands.  In this  paper  we focus  our  study on  the properties  of the  circumnuclear emission  seen in  the  Chandra soft X-ray images in  the complete sub-sample of nearby  ($z<0.3$) 3CR radio galaxies with High Excitation (HEGs) or Broad (BLO) emission lines.  The first case reported in  the literature of extended soft X-ray emission (below  $\\sim$2 keV)  associated with  an  AGN is  the Seyfert  galaxy NGC4151 (\\citealt{elvis83}) observed  with {\\it Einstein}. Since then, it has  been recognized that a  few bright, nearby  Seyfert 2 galaxies are associated with soft X-ray emission  extending over $\\sim$1 kpc and matching very   closely   the   morphology    of   the   optical   NLR   (e.g., \\citealt{elvis90}). This result was strengthened by later observations by the {\\it   Einstein}  and   ROSAT   satellites  (e.g.,   NGC~1068: \\citealt{wilson92},   NGC~2992:  \\citealt{elvis90},  NGC~2110: \\citealt{weaver95}). This was not a totally unexpected result, because a hot  medium in pressure equilibrium with  the NLR is thought to  be necessary to prevent the  optical emitting line clouds  from evaporating  \\citep{krolik84}. In  fact initially the preferred explanation of these extended soft X-ray regions was an outflow of hot, collisionally ionized gas, that confines the narrow-line clouds (\\citealt{wilson92}, \\citealt{elvis90}, \\citealt{weaver95}), although  other emission mechanisms, such as scattering of nuclear light, could not be ruled out.  Recently, with the advent of a new generation of higher resolution and sensitivity  X-ray telescopes,  such as  Chandra and  XMM-Newton, more detailed comparison  on sub-arcsec  scales between the  X-ray, optical, and  radio  emission  has   been  possible.  Chandra  images  at  high resolution and  XMM spectroscopic observations, combined  with HST and VLA images, have  been used to investigate a  larger number of bright, nearby    AGN,    mostly     Seyfert    galaxies    (e.g.,    NGC~2110 \\citealt{evans06}). \\citealt{bianchi06} studied a sample of 8 Seyferts and found extended soft X-ray emission co-spatial with the NLR for all of their  sources. High  resolution spectroscopic observations  of the extended X-ray emitting regions performed with Chandra/HETGS, possible in   a   few   cases   (e.g.   NGC~1068   \\citealt{ogle03},   NGC~4151 \\citealt{wang11}) reveals that the soft  X-ray part of the spectrum is dominated by emission  lines mainly from He- and  H-like K transitions of light metals similar to those observed in their nuclei (e.g. MRK 3: \\citealt{sako00}, Circinus:  \\citealt{sambruna01}). X-ray observations made   with  the  Reflection   Grating  Spectrometers   (RGS)  onboard XMM-Newton have lower resolution (15$\\arcsec$), thus encompassing both the nucleus and the soft  X-ray extended region. The resulting spectra therefore represent intensity-weighted conditions over kpc scales, but they  are consistent with  an ensemble of several  narrow emission lines (e.g.  \\citealt{kinkhabwala02}, \\citealt{guainazzi07}, NGC~5252: \\citealt{dadina10}).  Nevertheless,  although  in  Seyfert  2  galaxies  the  most  probable explanation is  that the gas is  photoionized by the AGN,  there is no general consensus  about the dominant  process that produces  the soft X-ray  extended emission  in other  classes of  galaxies.  In  fact, a variety of  different mechanisms could  be considered, such as:  a hot component  of a multiphase  interstellar medium,  hot gas  shocked and evacuated  by an  outflow or  a jet,  or an  outflow  of collisionally ionized  gas escaping  from  the nucleus,  driven  by radiation,  that interacts  with  the  NLR  clouds.  In an  \u201doutflow\u201d  model,  the  NLR kinematics  are dominated  by radiation  and/or wind  pressure driving clouds outwards  from the  nucleus \\citep{crenshaw00}  and indeed complex motions in the NLR are often inferred.  The literature on the X-ray properties of the NLR in radio galaxies is more limited  than for  Seyferts. At high  energies (between 2  and 10 keV)  radio galaxies  show a  compact nuclear  component and,  in some cases, collimated structures co-spatial  with the radio emission. Only recently extended regions in soft  X-rays, similar to that observed in Seyfert   galaxies,   have   been   discovered   (e.g.,   in   3C~171: \\citealt{hardcastle2010},    3C~33:    \\citealt{torresi09},    3C~305: \\citealt{massaro09}  and \\citealt{hardcastle12},  and  PKS 1138\\--262: \\citealt{carilli02}). In all cases the soft X-ray emitting regions are closely spatially  related to the  optical emission lines,  similar to what is observed for radio quiet AGN. Also, analysis of the nuclei performed with  high or  medium resolution spectroscopic  data reveals many emission lines  in the soft band, as  observed in the radio-quiet AGN (e.g.  3C~445: \\citealt{sambruna2007} et  \\citealt{grandi2007}, 3C~33: \\citealt{torresi09}, 3C~234: \\citealt{piconcelli2008}).  In this work we perform a complete analysis of the properties of X-ray emission in  the soft band (0.5-2  keV) for 113 3CR radio galaxies at $z<0.3$  (from \\citealt{spinrad1985}).  A Chandra snapshot survey has been recently completed  to provide us with observations of all 3CR sources not already covered by other programs. Now we can take advantage   of  a   complete   database  of   high  resolution   X-ray observations. In  this paper we  focus on  the extended circumnuclear emission of the galaxies  classified as HEG and BLO. Our aim  is to  establish whether  extended  emission is  present, and  to explore the relationship with the overall multi-band properties of the galaxy.  The multi-wavelength data we will use are mainly from the ground based spectroscopic              survey             presented             by \\citet{buttiglione09,buttiglione10,buttiglione11},    leading   to   a classification into the three  main classes: broad line objects (BLO), low and high excitation galaxies (LEG and HEG).  Furthermore, emission line imaging  surveys of the 3CR  sources have been  carried out using both ground-based telescopes  \\citep{mccarthy95,baum88} and the Hubble Space Telescope \\citep{privon08,tremblay09}.  We adopted the following cosmological parameters: H$_0$=71 km s$^{-1}$ Mpc$^{-1}$,          $\\Omega_M$=0.27,          $\\Omega_{\\Lambda}$=0.73 \\citep{spergel03}. ",
        "conclusions": ""
    },
    "1208/1208.0808_arXiv.txt": {
        "abstract": "The high-energy-peaked BL Lac H\\,2356-309 (z=0.165) was detected by HESS at very high energies (VHE, $\\gtrsim100$ GeV) with relatively high significance in the years 2004-2007, allowing a good determination of its gamma-ray spectrum. After correction for the interaction with the diffuse extragalactic background light (EBL), the VHE spectrum is  flat ($\\Gamma\\sim1.9-2$) over a decade in energy, locating the gamma-ray peak around or above 0.6-1 TeV. This is remarkably at odds with the interpretation and modeling provided by HESS, which do not correspond to the source properties and  can be excluded with high confidence. The overall GeV-to-TeV characteristics of H\\,2356-309 seem intermediate between the ``TeV-peaked\" (Fermi-faint) and ``100 GeV-peaked\" (Fermi-bright) BL Lac objects, and difficult to reconcile with the shape of the synchrotron emission in a single-zone SSC scenario. ",
        "introduction": "The BL Lac object H\\,2356-309 (z=0.165) is a High-frequency-peaked BL Lac (HBL), also called High-Synchrotron-Peaked blazar (HSP, \\citep{2lac}). It is usually bright in the X-ray band, and during BeppoSAX observations in 1998 it presented a synchrotron peak in the spectral energy distribution (SED) above few keV, which is the defining characteristic of the so-called ``extreme BL Lacs\" \\citep{extreme}.  At VHE, H\\,2356-309  was detected for the first time in 2004  by HESS \\citep{nature,hess2356}. The VHE spectrum was found to be significantly harder than expected from a source at that redshift, considering the softening effects of $\\gamma$-$\\gamma$ absorption on  the diffuse extragalactic  background light (EBL), in the detected energy range. Together with other even harder sources, this fact lead to the discovery of a low intensity of the EBL in the component produced by the direct starlight \\citep{nature}.  Since 2004, H\\,2356-309 has been monitored by HESS for several years, reaching a total detection of $\\sim$13 $\\sigma$ in the timespan 2004-2007, with  an average flux at the level of $\\sim$1.6\\% of the Crab flux (above 240 GeV). In this epoch, three simultaneous multi-wavelength campaigns were performed, in 2004 with RXTE and 2005 with XMM-Newton \\citep{hess2356,icrc08}, which allowed the characterization of its SED. The results of these campaigns, the VHE monitoring and a synchrotron self-Compton (SSC) modeling of the data were published by the HESS Collaboration in Abramowski et al. 2010 \\citep{2356mwl}.  In that paper, however, while all the data analysis is correct, the HESS Collaboration has apparently misinterpreted the VHE data, providing a wrong assessment of the intrinsic VHE properties of the source. It is also not consistent with previous publications on the same results. This contribution presents the arguments against the interpretation in \\citep{2356mwl} and tries to provide a more accurate description of the actual gamma-ray properties (from the GeV to the TeV band) of this BL Lac object. ",
        "conclusions": ""
    },
    "1208/1208.2499_arXiv.txt": {
        "abstract": "The phase transition from normal hadronic matter to quark matter in neutron stars (NS) could give rise to several interesting phenomena. Compact stars can have such exotic states up to their surface (called strange stars (SS)) or they can have quark core surrounded by hadronic matter, known as hybrid stars (HS). As the state of matter of the resultant SS/HS is different from the initial hadronic matter, their masses also differ. Therefore, such conversion leads to huge energy release, sometimes of the order of $10^{53}$ ergs. In the present work we study the qualitative energy released by such conversion. Recent observations reveal huge surface magnetic field in certain stars, termed magnetars. Such huge magnetic fields can modify the equations of state (EOS) of the matter describing the star. Therefore, the mass of magnetars are different from normal NS. The energy released during the conversion process from neutron magnetar (NM) to strange magnetar/hybrid magnetar (SS/HS) is different from normal NS to SS/HS conversion. In this work we calculate the energy release during the phase transition in magnetars. The energy released during NS to SS/HS conversion exceeds the energy released during NM to SM/HM conversion. The energy released during the conversion of NS to SS is always of the order of $10^{53}$ ergs. The amount of energy released during such conversion can only be compared to the energy observed during the gamma ray bursts (GRB). The energy liberated during NM to HM conversion is few times lesser, and is not likely to power GRB at cosmological distances. However, the magnetars are more likely to lose their energy from the magnetic poles and can produce giant flares, which are usually associated with magnetars. ",
        "introduction": "GRB are cosmic gamma ray emission distributed isotropically originating at extra galactic distances. They are inferred to have energy of the order of $10^{53}$ ergs. If the beaming and the $\\gamma$-ray production of the GRB are quite high, then the central engine which powers such conversion can release energy to power GRB at cosmological distances. There are various models describing the central engine for GRB. Mergers of two NS or a NS and a black hole in a binary \\cite{pacz} being some popular models. However, recent calculation of Janka \\& Ruffert \\cite{janka} had shown that the energy released in such neutrino-antineutrino annihilation is much smaller to account for the GRB, thought to be occurring at large cosmological distances. Therefore, the central engine which powers such huge energies in GRB still remains unclear.  On the other hand, the conversion of a NS to a QS liberates huge amount of energy, simply because the star mass changes during the conversion. Therefore, the mass difference of NS and QS manifest itself in the form of energy release. The first calculation of the energy release by such conversion was proposed by Alcock et al. \\cite{alcock} and Olinto \\cite{olinto}. More detailed similar calculation for the conversion of NS to SS/HS was done by Cheng \\& Dai \\cite{cheng}, Ma \\& Xie \\cite{ma} and subsequently by many others \\cite{bhat1,bhat2,drago,neibergal,herzog,mallick1,pagliara}. Almost all of them found the energy released to be greater than $10^{51}$ergs, and connected them with the observed gamma ray bursts (GRB). However, new observation \\cite{kulkarni98,kulkarni99} reveals the GRB to occur at cosmological distances. Therefore, energies of the order of $10^{51}$ ergs are too low to power GRB at such huge distances. A systematic calculation done by Bombaci \\& Datta \\cite{bombaci}, using various EOS, estimated the energy released during conversion to be of the order of $10^{53}$ergs. Their main assumption was that the initial NS baryonic mass (BM) and the final NS BM is the same, the BM is conserved during such conversion. Their estimation of the energy release can power GRB to such large distances and was in agreement with the observed GRB. However, during the conversion there can be ejection of matter from the outer layers of the star due to several shock bounce. The matter ejected is usually assumed to be low, but Fryer \\& Woosley \\cite{woosley} showed that the matter ejected can be as high as $0.2$ solar mass. For such a scenario the assumption of baryonic mass conservation does no longer hold.  The above calculation of energy release depends on the fact that some of the compact stars are really exotic and quark stars (QS either SS or HS) do really exist. However, the recent measurement \\cite{Demorest10,new2m} found the pulsars to have a very high mass, $M\\sim 2 $M$_\\odot$. Pulsars with such high masses, for the first time, has imposed strict constraints in determining the equations of state (EOS) of matter describing compact stars. The EOS of quark matter usually has strangeness \\cite{itoh,bodmer,witten} and therefore, provides an additional degree of freedom. This extra degree of freedom softens the EOS, that is, the reduction of pressure for a given energy density. As a result, it becomes difficult for the quark EOS to generate stars with such high mass which satisfies the mass limit set by new measured heavy pulsars. However, new calculations have found that the effects due to strong interaction, such as one-gluon exchange or color-superconductivity can make the EOS stiffer and increase the maximum mass limit of SS and HS \\cite{Ruester04,Horvath04,Alford07,Fischer10,Kurkela10a, Kurkela10b}. Ozel \\cite{Ozel10} and Lattimer \\cite{Lattimer10} were the first to study the implications of the new mass limit for SS and HS within the MIT bag model. Therefore, the conversion of NS to QS (SS/HS) is still a viable scenario of astrophysical phase transition (PT).  The only other astrophysical events which can come close to the energy budget of GRB are the recently observed giant flares. Three giant flares, SGR 0526-66, SGR 1900+14 and SGR 1806-20 \\cite{flare} have been detected so far. The huge amount of energy in these flares can be explained by the presence of a strong surface magnetic field, whose strength is estimated to be larger than $10^{14}$G, and such stars are termed magnetars. Simple calculation suggests that the magnetic fields in the star interior can be few order higher ($\\sim 10^{18}$G). Theoretical calculation also suggests that such strong magnetic field can effects the gross structure of NS. It effects the structure of a NS through the metric describing the star \\cite{bocquet,cardall} or by modifying the EOS of the matter through the Landau quantization of charged particles \\cite{chakrabarty97,yuan,broderick,chen,wei}. As the EOS gets modified, the mass of magnetar differs from a normal pulsar. This would eventually effect the energy released during the conversion of a NM to quark magnetar (QM). Therefore, the energy released during the PT of normal pulsars would be quite different from the energy released during the PT of magnetars.  In the present work, we calculate the energy release during the conversion of NS to SS/HS. Previous calculation does not include the idea of matter ejection during the conversion process. In our calculation we assume that the baryonic mass conservation to be true, only after taking into account the mass of the matter ejected. Finally, we would study the energy released during the conversion of a NM to SM/HM. In the next section (Section II), we give the outline of the calculations involved in the energy release. In section III, we discuss the EOS. We discuss the effect of magnetic field on the EOS in section IV, and the HS is discussed in section V. The results are discussed in section VI. Finally, in section VII, we summarize our results and draw conclusion from them. ",
        "conclusions": "In this work, we have calculated the energy released during the conversion of NS to QS (SS/HS). The total energy released is the sum of the gravitational and internal energy of conversion. Our first assumption was that the conversion starts with a sudden spin down of the star, resulting in a huge density fluctuation at the core, initiating the phase transition. We also assume that some matter is ejected from the outer layers of the star due to several shock bounce, and therefore, the change in baryonic mass. The conversion may continue up to the surface or may die out after some distance. This depends on the energy difference between the matter phases at the centre of the star and also on the initial density and spin fluctuation of the star. The final star may be a SS or a HS. For the NS we have considered relativistic mean field EOS model of hadronic matter. For the quark matter (SS/HS), we have considered simple MIT bag model. We construct the HS based on Glendenning construction. In our calculation we have used two different parametrization for the hadronic EOS, and had regulated the quark EOS by changing the bag constant.  First we have shown the energy liberated during the conversion of NS to QS, with no matter ejection. With a fixed baryonic mass, the energy liberated is obtained from the difference in the gravitational mass of the initial NS and final QS. For such a case the liberated energy is always close to $10^{53}$ ergs.  Next, we had shown the energy liberated during the conversion of NS to QS, with matter ejection from the surface layer of the NS. The amount of matter ejected is taken to be $0.2 M_{\\odot}$ of $M_B$. Therefore, the resulting baryonic mass denoted as $M_{BE}$ is $(M_B -0.2) M_{\\odot}$. Due to matter ejection as the $M_B$ decreases the gravitational mass and proper mass of the NS and QS also decreases. The energy liberated during the conversion of NS to QS with matter ejection is always less than that of the energy liberated during the conversion of NS to QS with no matter ejection. We have also studied the conversion of NS to SS/HS having high surface magnetic field (observed magnetars). We have denoted them as NM and QM (SM/HM). The energy liberated during the conversion of magnetars is less than that for normal pulsars.  One thing that clearly points out is that the conversion of NS to HS liberates energy few times less than that for the conversion of NS to SS. Therefore, it would be difficult for them to power GRB from cosmological distances. For the Max HS the energy liberated is the lowest and sometimes it is zero, which indicates the PT there occurs without any observable signal.  The energy liberated during the conversion of NS to QS is of $10^{53}$ ergs, and it can account for the cosmological origin of the GRB if the Lorentz factor is high. The energy released during the conversion of NM to QM is of the order of $10^{52}$ ergs, and are therefore difficult to power GRB at cosmological distances. However, the magnetars would release energy more efficiently from their magnetic poles than the normal NS. Therefore, in magnetars the energy liberated from the star with large Lorentz number and can account for the giant flares activity.  The detailed picture of the conversion mechanism and the central engine for the GRB is still not well understood. However, the above described processes are thought to be responsible for the energetic of the GRB. In this paper, we have concentrated more on the difference in gravitational binding energy due to phase transition which is accompanied by matter ejection from the outer layers. We have seen the difference in the energy release for normal pulsars and magnetars. The amount of matter ejected, the fraction of energy which goes towards heating of the star and the fraction spent on maintaining the conversion would ultimately determine the fate of energy release. However, if a significant amount of energy released goes into the energetic of electron-positron pairs production, it can account for the GRB energy.  Assuming the baryonic mass of the star to remain unaltered, Bombaci and Dutta \\cite{bombaci} performed their calculation and found that the NS has always greater mass than the final SS. Neutrino-antineutrino pair annihilation to electron-positron pairs deposit energy of the order of $10^{49}$ergs \\cite{cheng}. Neutron proton scattering by neutrinos inside the dense star deposit energy at least two to three order further higher. Due to the shock of the phase transition some matter near the stellar surface may be ejected and would further be accelerated by electron-positron pairs and may eventually oscillate. This may give rise to high Lorentz factor and high luminosity needed for GRB. These two are the most efficient processes to account for the energetic of the GRB. As the actual process takes place at cosmological distances, we do not have a better understanding of such phenomena. As there are no other source of external energy available to the star, other than the rotational energy, a huge amount of energy may be used in the actual conversion. We can only conclude that if a small fraction of the energy of the conversion is released it may manifest itself at least in the form of giant flares, which are usually associated with the magnetars. However, if a substantial amount of energy is released, it may account for the energies observed during the GRB. Therefore, more detailed studies in theoretical and observational front is still needed for the pulsars and magnetars to have a better understanding of the energetic of the conversion.  On the other hand, with new observation of $2$ solar mass pulsar gives very strong constraints for the EOS describing compact stars. The EOS used in our calculation cannot satisfy the new mass limit. However, this can be remedied by considering new prescriptions for both hadronic and quark matter EOS. With the improved EOS, the basic results of our calculation does not have any qualitative change but only small quantitative change. With new techniques of precise mass measurement, the physics of pulsars is entering a new phase. As the exact nature of strong interactions is still far from being settled, the quark EOS being very model dependent, the maximum mass of NS and QS calculated from theoretical consideration are going to evolve further in future. To get a clear picture of energy release and its detailed mechanism we need more detailed microscopic studies and this work is the first step towards that direction."
    },
    "1208/1208.5046_arXiv.txt": {
        "abstract": "The usual nuclear recoil energy reconstruction employed by liquid xenon dark matter search experiments relies only on the primary scintillation photon signal. Energy reconstruction based on both the photon and electron signals yields a more accurate representation of search results. For a dark matter particle mass $m_{\\chi}\\sim10$~GeV, a nuclear recoil from a scattering event is more likely to be observed in the lower left corner of the typical search box, rather than near the nuclear recoil calibration centroid. In this region of the search box, the actual nuclear recoil energies are smaller than the usual energy scale suggests, by about a factor $\\times2$. Recent search results from the XENON100 experiment are discussed in light of these considerations. ",
        "introduction": " ",
        "conclusions": "Nuclear recoil energy reconstruction in liquid xenon has historically been defined by ${E}_{nr}~=~\\langle \\mbox{S1} \\rangle /(L_y\\mathcal{L}_{eff})(S_e/S_n)$, where $\\mathcal{L}_{eff}$ is the relative scintillation efficiency of nuclear recoils (also referred to as the effective Lindhard factor). A monoenergetic gamma (usually 122~keV), with a measured light yield $L_y$ given in photoelectrons/keV, is used as standard candle. An electric field applied across the xenon target allows the measurement of the electron signal, and it also quenches the scintillation signal. The quenching of the gamma is $S_e \\sim 0.5$, while nuclear recoils are barely quenched ($S_n\\sim0.95$). Both $L_y$ and $S_e$ vary significantly with the electric field strength.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=0.48\\textwidth]{fig1.pdf} \\vskip -0.1cm \\caption{{\\bf (upper panel)} Simulated nuclear recoil calibration data from a hypothetical detector which is similar to the XENON100 detector during its 2009 commissioning run. The dashed curves, described in the text, define three sides of the signal search box. The inset contours indicate $\\langle {E}_{nr} \\rangle$. Electron recoils from electromagnetic background have $\\langle y \\rangle=0$. The energy scale used in \\cite{:2012nq} is indicated along the top, and the two events reported therein are reproduced for comparison. {\\bf (lower panel)} S1 acceptance of the hypothetical detector (solid), and the XENON100 detector (dashed).} \\vskip -0.5cm \\label{fig:energyContours} \\end{center} \\end{figure}  Dedicated, direct measurements of the relative scintillation efficiency $\\mathcal{L}_{eff}$ are plentiful \\cite{Manzur:2009hp,Aprile:2008rc,Plante:2011hw}. These measurements report, for a given ${E}_{nr}$, the average number of recorded photoelectrons $\\langle \\mbox{S1} \\rangle$. In contrast, dark matter search experiments report a measured S1 (and S2) corresponding to each event, and wish to know the most likely nuclear recoil energy $\\langle {E}_{nr} \\rangle$ associated with that event. For events near the nuclear recoil band centroid (in $y$), these two operations commute. But in many cases, such as the recently reported two events in the XENON100 search region \\cite{:2012nq}, they do not.  It appears that a nuclear recoil energy reconstruction based directly on Lindhard theory is possible in liquid xenon \\cite{Sorensen:2011bd}. The method employed in this work is very nearly equivalent to that, but framed so as to make an explicit connection with the historical method, and with the experimentally measured quantities.  Figure \\ref{fig:energyContours} shows the results of a simulation of the nuclear recoil response of a hypothetical liquid xenon detector. The results are plotted in the customary variables, though other choices are possible \\cite{Arisaka:2012ce}, and may be useful considering the relative size of $\\alpha_1$ and $\\alpha_2$. The simulation method is described in \\cite{Sorensen:2010hq},  and reproduces the relevant binomial and Poisson statistical processes. It has been shown to provide a hi-fidelity reproduction of XENON10 nuclear recoil calibration data. The detector-specific details were obtained from \\cite{Aprile:2010um,Aprile:2011dd}. The hypothetical detector is therefore expected to exhibit a response similar to the XENON100 detector during its initial phase of operation. Uncertainties are discussed in a separate section. In this work, we took as inputs the nuclear recoil centroid reported in Fig. 3 of \\cite{Aprile:2010um}, and the central $\\mathcal{L}_{eff}$ curve from \\cite{Aprile:2011hi} (shown in Fig. \\ref{fig:energyScale}, lower panel, solid curve). The simulated data are shown after subtracting the electron recoil centroid ($\\mu_{ER}$), which was also taken from \\cite{Aprile:2010um}.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=0.48\\textwidth]{fig2.pdf} \\vskip -0.1cm \\caption{Nuclear recoil energy calibration data for liquid xenon. The simulation assumes the solid curves. Events with $E_{nr}<2$~keV were not simulated. $\\mathcal{L}_{eff}$ data reproduced from \\cite{Manzur:2009hp} (stars), \\cite{Aprile:2008rc} (diamonds) and \\cite{Plante:2011hw} (squares). Additional $\\mathcal{Q}_y$ data from \\cite{Aprile:2006kx}. } \\vskip -0.5cm \\label{fig:energyScale} \\end{center} \\end{figure}  The simulation has only a single free parameter, the ionization yield $\\mathcal{Q}_y\\equiv\\mbox{S2}/{E}_{nr}$. It was allowed to float until the simulation band centroid matched the nuclear recoil centroid from data \\cite{Aprile:2010um}. The agreement is very good, within $1\\sigma$ of the statistical uncertainty on the mean, above $\\mbox{S1}=3$. Below $\\mbox{S1}=3$, the agreement is within $2\\sigma$. The $\\mathcal{Q}_y$ curve so obtained is shown in Fig. \\ref{fig:energyScale} (solid curve). This does not guarantee that either $\\mathcal{L}_{eff}$ or $\\mathcal{Q}_y$ are correct in absolute terms, but rather, as drawn, are self-consistent with nuclear recoil band data. We note that the lower (dashed) $\\mathcal{Q}_y$ curve is very consistent with the NEST model \\cite{Szydagis:2011tk}.  In addition to the simulated nuclear recoil data, Fig. \\ref{fig:energyContours} shows the centroid $\\mu_{NR}$ (dashed, green) and $\\mu_{NR}-3\\sigma_{NR}$ (dashed, black). Also shown are the ``software'' $\\mbox{S1}>3$ (dashed) and $\\mbox{S2}>150$ (stippled) thresholds, as in \\cite{:2012nq}. The dashed curves define three walls of the dark matter search box for our hypothetical detector. In a similar (actual) search box, the XENON100 Collaboration recently reported the observation of two events, which are reproduced here. The stated event energies are 7.1 (circle) and 7.8 keV (star), and this appears very reasonable according to the  ${E}_{nr}\\propto \\mbox{S1}$ scale given along the top axis.  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.48\\textwidth]{fig3.pdf} \\vskip -0.1cm \\caption{The probability to observe a fluctuation equal to or greater than the two events shown in Fig. \\ref{fig:energyContours}, as a function of nuclear recoil energy. The abcissa is the actual simulated energy (not the contours in Fig. \\ref{fig:energyContours}), and the markers correspond to the two events.} \\vskip -0.5cm \\label{fig:energyProbability} \\end{center} \\end{figure}  We used the simulated events to find contours of  $\\langle {E}_{nr} \\rangle$, as shown in Fig. \\ref{fig:energyContours} (solid curves, with corresponding energy in keV). Because a downward fluctuation in S1 is accompanied by an upward fluctuation in $y$, the contours follow the S2 expectation value for each energy. This is significantly different from the cartesian expectation implied by ${E}_{nr}\\propto \\mbox{S1}$ (notice that this scale is most correct near the calibration centroid, near $y\\approx-0.4$). Figure \\ref{fig:energyProbability} shows the probability that a nuclear recoil of energy $E_{nr}$ resulted in either of the two observed events. Only simulated events which produced a measurable S1 and S2 signal were considered.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.98\\textwidth]{fig4.pdf} \\vskip -0.1cm \\caption{The expected distribution of nuclear recoil events, for several values of $m_{\\chi}$. In order to clearly show the distribution, the following cross sections $\\sigma_n$ were assumed: $5\\times10^{-39}$, $1\\times10^{-40}$, $1\\times10^{-41}$ and $1\\times10^{-41}$~cm$^2$.} \\vskip -0.5cm \\label{fig:examples} \\end{center} \\end{figure*}  The discussion up to this point has assumed a nuclear recoil spectrum corresponding to an americium-berylium neutron source, as is frequently used to calibrate the nuclear recoil response of liquid xenon detectors. In Fig. \\ref{fig:examples} we show the expected distributions of events for several dark matter masses $m_{\\chi}$. The spectral distributions were calculated assuming the same astrophysical parameters described in \\cite{:2012nq}. This clearly shows, particularly for $m_{\\chi}\\lesssim10$~GeV, that the $y$ coordinate also carries spectral information. The point is perhaps obvious from the definition of $y$, but it has been neglected in previous work. %  \\mysection{{\\it Uncertainties}} \\label{sec:uncert} The ordinary statistical processes modeled by the simulation lead to  non-Gaussian tails in the distribution of $y$, particularly for $\\mbox{S1}\\lesssim10$. The extent of the tails are roughly indicated by the $\\langle {E}_{nr} \\rangle$ contours in Fig. \\ref{fig:energyContours}. In Fig. \\ref{fig:energyScale} (upper panel), we indicate the approximate range of $\\mathcal{Q}_y$ (dashed curves) which adequately reproduce the nuclear recoil band data centroid in \\cite{Aprile:2010um}, given $\\mathcal{L}_{eff}$ as in the lower panel (solid curve). The lower dashed curve is the model prediction employed in \\cite{Angle:2011th}; $\\mathcal{Q}_y$ below this do not appear reasonable. It can be seen that the uncertainty in $\\mathcal{Q}_y$ is most significant for ${E}_{nr}\\lesssim10$~keV. This is due primarily to the non-Gaussian tails in $y$, and how the mean value of the distribution is determined (our method may be slightly different from what was used in \\cite{Aprile:2010um}). As a result, one may expect a systematic shift in the $\\langle {E}_{nr} \\rangle$ contours in Fig. \\ref{fig:energyContours}, of as much as $\\Delta y = ^{+0.15}_{-0.06}$ for ${E}_{nr}\\lesssim10$~keV. %  The S1 acceptance of our hypothetical detector response is different from  \\cite{:2012nq}, as shown in Fig. \\ref{fig:energyContours} (lower panel). The statistical methods we employ make it easy to adjust the predicted S1 acceptance to match that of \\cite{:2012nq}. This has a systematic effect on the best-fit $\\mathcal{Q}_y$, as pointed out in \\cite{Aprile:2012he}, and leads to a maximum displacement $\\Delta y = +0.02$ of the $\\langle {E}_{nr} \\rangle$ contours. This effect causes the largest shift in the range $3-5$~keV.  Of course, there is also uncertainty in $\\mathcal{L}_{eff}$ itself. In this work, we have taken as a prior the $\\mathcal{L}_{eff}$ favored by the XENON100 Collaboration. An equally plausible choice would be \\cite{Manzur:2009hp} (shown in Fig. \\ref{fig:energyScale}, stars). A 20\\% smaller $\\mathcal{L}_{eff}$ would require a systematic decrease in $\\mathcal{Q}_y$ of about $8-12\\%$. This would lead to an essentially uniform $-5\\%$ systematic shift in the $\\langle {E}_{nr} \\rangle$ contours, as can be verified analytically.  It is also notable that the S1 response reported in \\cite{:2012nq} has improved by $\\sim4\\%$ and the S2 response has improved by $\\sim14\\%$, relative to \\cite{Aprile:2010um} (and hence relative to our hypothetical detector). This leads to a uniform $+4\\%$ systematic shift in the $\\langle {E}_{nr} \\rangle$ contours. We mention in passing that the width of the simulated band (in $y$) appears in good agreement with previously reported results \\cite{Aprile:2010um,Aprile:2011hi}, but slightly narrower than \\cite{:2012nq}.  In spite of these uncertainties, the fundamental profile of the $\\langle {E}_{nr} \\rangle$ contours remains as shown in Fig. \\ref{fig:energyContours}.  \\mysection{{\\it Conclusions}} \\label{sec:conclus} Full consideration of the energy information carried by both the S1 and S2 signals indicates that the energies of the two events observed by the XENON100 detector would be reconstructed at $2.9\\pm0.5$~keV and $3.6\\pm0.6$~keV in our hypothetical detector, subject to a systematic uncertainty of $^{+0.3}_{-1.0}$~keV. This is about a factor $\\times2$ smaller than what one obtains from ${E}_{nr}~\\propto~\\mbox{S1}$. Interestingly, the approximate location of these events (near the lower left corner of a typical search box, rather than near the nuclear recoil centroid) is what one would expect for elastic scattering of low-mass ($m_{\\chi}\\lesssim10$~GeV) dark matter. This observation could be important from a phenomenological perspective ({\\it e.g.} \\cite{Hooper:2012ft,Hooper:2012cw}).  It is evident from Fig. \\ref{fig:examples} is that the acceptance of the search box (as a function of S1) depends on $m_{\\chi}$. For the four example $m_{\\chi}$ values, the fraction of events below the calibration centroid in the range $3\\leq\\mbox{S1}<10$ are 1.00, 0.99, 0.56 and 0.42. This is a result of the shape of the calculated dark matter energy spectrum.  The fraction is about 0.53 for the calibration data, slightly above 0.50 due to the non-Gaussian tails. Energy reconstruction based on $E_{nr} \\propto \\mbox{S1}$ leads to the assumption that the acceptance of the search box is always given by the calibration data. Relative to previously reported results, this should tend to strengthen the sensitivity to particle masses $m_{\\chi}\\lesssim35$~GeV, and weaken the sensitivity to larger masses. From an experimental point of view, it is interesting that for $m_{\\chi}\\lesssim10$~GeV, the electromagnetic background population with $\\langle y \\rangle = 0$ is essentially irrelevant. %  \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.48\\textwidth]{fig5.pdf} \\vskip -0.1cm \\caption{The region of parameter space in which an outcome similar to the two events shown in Fig. \\ref{fig:energyContours} is at least 10\\% likely (crosses), and the 90\\% CL exclusion limits from XENON100 \\cite{:2012nq} (thick curve) and XENON10 \\cite{Angle:2011th} (thin curve).  {\\bf (Inset)} the result of a simulated experiment, for $\\sigma_n=3\\times10^{-43}~\\mbox{cm}^2$ and $m_{\\chi}=7$~GeV. The two simulated events are shown as large circles; regions (a) and (b) are described in the text; markings and axes are as described in Fig. \\ref{fig:energyContours}. } \\vskip -0.5cm \\label{fig:limits} \\end{center} \\end{figure}  In Fig. \\ref{fig:limits} we show values of $(m_{\\chi} , \\sigma_n)$, in which outcomes similar to that recently observed by  XENON100 appear at least 10\\% likely (crosses) in our hypothetical detector. The S1 acceptance is shown in Fig. \\ref{fig:energyContours} (solid curve), and a 34~kg~$\\times$~224.6 day exposure was assumed. Our definition of ``similar'' is two events observed in the search box below the $\\langle {E}_{nr} \\rangle = 5$~keV contour, and no events above it. The search box is shown in Fig. \\ref{fig:energyContours}, and Fig. \\ref{fig:limits} (inset), bounded by dashed lines. Additionally, to account for the uncertainties discussed above, we allow (a) the possibility that one of the two events appears in the region below the search box, but above the S2 threshold, and (b) the possibility that one of the two events falls in the region S1$<3$. These regions are indicated in Fig. \\ref{fig:limits} (inset). Uncertainties would propagate into Fig. \\ref{fig:limits} as follows: a larger S1 acceptance ({\\it e.g.} Fig. \\ref{fig:energyContours} dashed curve) would tend to push the $(m_{\\chi} , \\sigma_n)$ region to smaller $\\sigma_n$. A larger $\\mathcal{Q}_y$ would mean a smaller energy for events near the lower left corner of the search box, and would therefore tend to push the $(m_{\\chi} , \\sigma_n)$ region to smaller $m_{\\chi}$.  We do not suggest that these two events observed in \\cite{:2012nq} are due to the elastic scattering of dark matter; the background hypothesis is of a similar likelihood, and thus more compelling. However, we have shown that in considering detection scenarios, significant additional information is gained from an energy reconstruction based on both $n_e$ and $n_{\\gamma}$. Specifically, while a result similar to \\cite{:2012nq} is compatible with low-mass dark matter, it is highly unlikely to have arisen from dark matter with $m_{\\chi}\\gtrsim10$~GeV.      \\begin{center} {\\bf Acknowledgements} \\end{center} Thanks are due to Adam Bernstein, Rouven Essig, Rick Gaitskell, Jeremy Mardon, Neal Weiner and the XENON100 Collaboration, for suggesting improvements to the manuscript. This work was performed under the auspices of the U.S. Department of Energy by Lawrence Livermore National Laboratory under contract DE-AC52-07NA27344. Report number LLNL-TR-574054.  {\\it Note added.--} We thank the authors of \\cite{Davis:2012hn} for bringing their related work to our attention."
    },
    "1208/1208.4871_arXiv.txt": {
        "abstract": "We explore perturbations about a Friedmann-Robertson-Walker background in Chern-Simons gravity. At large momenta one of the two circularly polarized tensor modes becomes ghostlike.  We argue that nevertheless the theory does not exhibit classical runaway solutions, except possibly in the relativistic nonlinear regime.  However, the ghost modes cause the vacuum state to be quantum mechanically unstable, with a decay rate that is naively infinite.  The decay rate can be made finite only if one interprets the theory as an effective quantum field theory valid up to some momentum cutoff $\\Lambda$, which violates Lorentz invariance. By demanding that the energy density in photons created by vacuum decay over the lifetime of the Universe not violate observational bounds, we derive strong constraints on the two dimensional parameter space of the theory, consisting of the cutoff $\\Lambda$ and the Chern-Simons mass. ",
        "introduction": "General relativity has held up well to various tests over the years from experiments and astronomical observations \\cite{GRreview}, and is considered a pillar of standard cosmology. However, it is interesting to consider modifications to the theory, particularly in light of the observed acceleration of the Universe \\cite{Skordis}.  One useful approach is to think of Einstein gravity as an effective field theory, and consider higher order corrections to the Einstein-Hilbert action, either involving the metric alone or involving an additional posited scalar field. The goal then becomes to calculate the corrections to general relativity arising from these higher order terms, and using experiments to set bounds on the couplings parameterizing them.  One such extension to general relativity is Chern-Simons gravity \\cite{Jackiw,CSreview}, where one assumes the existence of a scalar field $\\vartheta$ coupled to gravity through a parity violating term. The theory is described by the action\\footnote{Throughout this paper we will restrict attention to the theory (\\ref{first}) in which the scalar field is dynamical; we will not consider the ``non-dynamical'' version of the theory in which the kinetic term for the scalar field is absent \\protect{\\cite{CSreview}}.  We note however that our derivation of the action (\\protect{\\ref{eq:quadratic}}) of the ghost graviton modes is valid for the non-dynamical theory, since the derivation does not involve any perturbations to the scalar field.} \\begin{equation}\\label{first} S = S_{EH} + S_{CS} + S_{\\vartheta} + S_{\\rm mat}, \\end{equation} where the various terms are respectively the Einstein-Hilbert term \\begin{subequations} \\begin{equation} S_{EH} = \\frac{1}{2}m^2_{\\rm p} \\int d^4x \\sqrt{-g} R, \\label{eq:EinsteinHilbert} \\end{equation} the Chern-Simons term \\begin{equation} S_{CS} = \\frac{1}{4} \\alpha \\int d^4x \\sqrt{-g} \\ \\vartheta \\ ^*RR , \\label{eq:CS}\\\\ \\end{equation} the scalar term \\begin{equation} S_{\\vartheta} = - \\frac{1}{2} \\int d^4x \\sqrt{-g} \\ [\\nabla_a \\vartheta \\nabla^a \\vartheta + 2V(\\vartheta)], \\label{eq:scalar} \\end{equation} \\end{subequations} and $S_{\\rm mat}$ describes any other matter present. In these expressions $m^2_{\\rm p} = (8\\pi G)^{-1}$ is the square of the reduced Planck mass, $g$ is the determinant of the metric, $R$ is the Ricci scalar and $\\alpha$ a coupling constant with dimensions of inverse mass.   Here and throughout we use units with $\\hbar = c = 1$.  Also $V(\\vartheta)$ is an arbitrary potential and the Pontryagin density is defined as \\begin{equation} ^*RR = \\frac{1}{2} \\epsilon^{cdef}R^a_{\\ bef}R^b_{\\ acd}, \\end{equation} where $\\epsilon^{cdef}$ is the four dimensional Levi-Civita tensor. For purposes of the present discussion, we will assume that matter is minimally coupled to the metric.   Consider now the dynamics of Chern-Simons gravity in perturbation theory about a Friedmann-Robertson-Walker (FRW) cosmological background. Now in the limit $\\vartheta =$ constant, the Chern-Simons term (\\ref{eq:CS}) reduces to a surface term and we recover Einstein's equations for the space-time dynamics.  Therefore, in the effective field theory that describes the perturbations, the operators that arise from the Chern-Simons term must be suppressed by a mass scale that is related to the derivative of the background scalar field.  This mass scale is called the Chern-Simons mass $m_{\\rm cs}$, and is defined in the FRW context by \\cite{solar} \\begin{equation} m_{\\rm cs} \\equiv \\frac{m^2_{\\rm p}}{\\alpha \\dot{\\vartheta}}, \\label{eq:CSscale} \\end{equation} where the dot denotes a derivative with respect to time.  General relativity is recovered in the limit $m_{\\rm cs} \\rightarrow \\infty$ for linearized tensor perturbations. Because it is the Chern-Simons mass that enters into equations describing observables for linear perturbations, we choose to constrain it, rather than the more fundamental coupling $\\alpha$ that appears in the action (\\ref{eq:CS}). Also, since the background cosmological solution $\\vartheta(t)$ need not be a linear function of time, the Chern-Simons mass will in general be a function of time or of redshift; we will focus in this paper on its value $m_{\\rm cs} = m_{\\rm cs}(t_0)$ today.     Past work constraining Chern-Simons gravity has utilized Solar System and binary pulsar tests of general relativity. Measurement of Lense-Thirring precession by LAGEOS give the bound \\cite{solar} \\begin{equation} m_{\\rm cs} \\gtrsim 2 \\times 10^{-13} \\, {\\rm eV}. \\end{equation} A bound $10^{11}$ times stronger has been claimed from binary pulsar studies \\cite{pulsar}, but the validity of this result has been questioned \\cite{pulsarcriticism}, and a corrected bound\\footnote{Although this bound was derived in the context of the non-dynamical Chern-Simons theory, without the kinetic term for the scalar field, it is also valid for the dynamical theory when the cosmological background solution $\\vartheta(\\eta)$ is nonzero, since the derivation is based on the modified gravitational Amp\\`ere equation which is valid in the dynamical theory \\cite{solar}.} from binary pulsars is \\cite{pulsarcriticism} \\begin{equation} m_{\\rm cs} \\gtrsim 5 \\times 10^{-10} \\, {\\rm eV}. \\label{fixed} \\end{equation}    In this paper we study the vacuum stability of Friedmann-Robertson-Walker (FRW) solutions in Chern-Simons gravity as a function of the Chern-Simons mass parameter. In Sec.\\ \\ref{sec:ghosts} we consider tensor perturbations to the FRW metric. We show that for spatial momenta above the Chern-Simons mass scale, one of the two polarization modes is ghostlike and can decay to radiation.  Requiring that the radiation produced over the lifetime of the Universe not exceed observational bounds allows us to constraint the parameters of the theory, which we do in Sec.\\ \\ref{sec:vacuumdecay}.  Finally in Section \\ref{sec:nonpert} we argue that the theory does not exhibit classical runaway solutions, despite the existence of ghostlike modes, except possibly in the relativistic nonlinear regime. ",
        "conclusions": ""
    },
    "1208/1208.1453_arXiv.txt": {
        "abstract": "A Chandra observation of the Large Magellanic Cloud supernova remnant DEM L241 reveals an interior unresolved source which is probably an accretion-powered binary.  The optical counterpart is an O5III(f) star making this a High-Mass X-ray Binary (HMXB) with orbital period likely to be of order tens of days. Emission from the remnant interior is thermal and spectral information is used to derive density and mass of the hot material. Elongation of the remnant is unusual and possible causes of this are discussed.  The precursor star probably had mass $> 25 M_{\\odot}$ ",
        "introduction": "In the Magellanic Clouds there are now 34 supernova remnants known to emit X-rays.  Ten of these have interior pulsar-wind nebulae (PWNe) or compact objects which also radiate in the X-ray band.  Although the sample is small, there is great diversity.  Several manifestations of neutron stars are represented and the present observation may be an example of yet another.  X-rays from the supernova remnant SNR 0535-67.5 in the H {\\footnotesize II} region DEM L241 (Davis et al 1976) were first detected in 1979 by Long et al (1981).   The supernova remnant was first identified by Mathewson et al (1985) and mapped using the optical [S {\\footnotesize II}] emission as seen in Figure \\ref{fig-opt-snr} which shows the remnant and the surrounding H {\\footnotesize II} emission. Since the initial Einstein observation there have been X-ray detections by ROSAT (Williams et al. 1999), XMM (Bamba et al. 2006), and now Chandra.  The XMM data provided the first detailed X-ray image and showed elongated diffuse emission filling the area outlined by [SII] filaments with a bright hard point source centered in the SE section.  Because the spectrum of this source was a power law, Bamba et al (2006) identified it as an unresolved PWN.  The Chandra observation was planned to resolve this object and to distinguish the expected point-like pulsar from surrounding diffuse emission.  The Chandra result indeed shows a clear point-like hard source but there is no sign of a PWN close to or surrounding the point source. ",
        "conclusions": "This is an interesting supernova remnant, larger and more elongated than most.  The absence of X-ray emission from the outer shock, the X-ray spectra, and the size all indicate an older remnant.  Because of its age and location close to a strong H {\\footnotesize II} region, radio emission is weak and undetectable.  It contains a compact object with an O-star optical counterpart which is unusual for a supernova remnant. No non-thermal X-rays, as might originate in a PWN, are detected from the immediate vicinity of this source and the source luminosity and spectrum are consistent with that expected from a HMXB.  The velocity of the O-star was observed to vary and a long period is suspected.  The diffuse X-ray spectrum from the remnant interior is enriched in O, Ne, and Mg.  This enrichment and the more-slowly evolving O-star companion imply that the supernova precursor star had a mass of $>25 M_{\\odot}$.   We note that the elongated envelope and the structure close to the compact source might have been formed by an SS 433-type pair of jets but there is no indication of these in the optical spectrum.  Support for this work was provided by the National Aeronautics and Space Administration through Chandra Award Number GO1-12094 issued by the Chandra X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory for and on behalf of the National Aeronautics Space Administration under contract NAS8-03060. D.L.F. acknowledges support from NASA through the Harriett G. Jenkins Pre-doctoral Fellowship Program, and from the Vanderbilt-University of Cape Town Partnership. We thank Paul Green for an interesting discussion and information about background quasars.  Sean Points supplied calibrated and aligned MCELS images. B. Furnish, J. Hood, C. McCarty, and T. Williams at Columbus State University helped with the X-ray spectral analysis."
    },
    "1208/1208.4276_arXiv.txt": {
        "abstract": "In supersymmetric models of warm inflation, the large temperature of the radiation bath produced by the dissipative motion of the inflaton field may induce a significant thermal abundance of potentially dangerous gravitinos. While previous discussions of this problem focused on gravitino production only at the end of warm inflation, similarly to conventional reheating scenarios, we study the full evolution of the gravitino abundance during and after inflation for simple monomial potentials, taking into account the enhanced gravitino and possibly gaugino masses due to supersymmetry breaking during inflation and the smooth transition into a radiation-dominated era. We find, on one hand, that the continuous thermal production increases the gravitino yield, although, on the other hand,  `freeze-out' occurs at temperatures much lower than previously estimated. Moreover, for sufficiently strong dissipation, which allows for sub-planckian inflaton values, the lower radiation temperature significantly alleviates and possibly solves the gravitino problem, with a baryon asymmetry being nevertheless produced through dissipative effects. Our analysis may also be relevant to standard reheating as an oscillating inflaton will also change the gravitino mass, potentially modifying the produced gravitino yield. ",
        "introduction": "Inflation \\cite{Guth1981a,Albrecht:1982wi,Linde:1981mu} has been incredibly successful in providing solutions to the problems of the standard cosmological model. It can set the initial conditions, which give rise to the high degree of flatness and homogeneity that we observe in the universe today. From particle physics motivated models, it not only yields a mechanism for accelerated expansion but also explains, through quantum fluctuations, the origin of the temperature anisotropies in the Cosmic Microwave Background and the seeds for the observed Large Scale Structure.  In the standard cold or isentropic inflation scenario, the early universe is dominated by the vacuum energy of a scalar field which is slowly rolling down its potential, resulting in a period of accelerated expansion. This occurs whilst its kinetic energy is negligible compared to the potential energy until, at some point, the potential typically steepens and the inflaton begins oscillating about the minimum of its potential. During the period of accelerated expansion, the inflaton is assumed to have negligible couplings to other fields in order to keep the potential flat enough for a sufficient number of e-folds of inflation ($\\sim$40-60) to occur and as a result the universe supercools. However, once it begins oscillating, there must be interactions that convert the vacuum energy into radiation in order to reheat the universe.  Hence, the inflaton cannot be an isolated system and, while the standard picture assumes that any interactions have a negligible effect on the dynamics of inflation and only become important during reheating, this need not be the case. In the alternative warm inflation paradigm, such interactions may in fact lead to dissipation of the inflaton's kinetic energy into light degrees of freedom, which in the simplest case may thermalise, resulting in the presence of a nearly-thermal bath concurrent with the accelerated expansion.  In the early universe, gravitinos can be abundantly produced, potentially leading to overclosure of the universe or spoiling the abundances of light elements predicted by the standard big bang nucleosynthesis model (BBN). While in cold inflation thermal production of gravitinos occurs only during the reheating phase, in warm inflation this is concurrent with inflation due to the presence of a thermal bath. In standard reheating, gravitino overproduction constrains the reheat temperature, i.e. the maximum temperature after inflation when the universe becomes radiation dominated. There is, however, a certain amount of tension in this case between having a large enough reheat temperature to allow for a thermal mechanism for baryogenesis, whilst keeping it low enough to avoid overproducing gravitinos. In warm inflation, on the other hand, this tension can be relieved, as a baryon asymmetry may in fact be produced at low temperatures through dissipative effects \\cite{Bastero-Gil2011}, potentially avoiding overproduction of gravitinos.  Gravitino production in warm inflation has been considered previously in \\cite{Taylor2001,Bastero-gila}, where it was assumed that the effective `reheat temperature' occurs when the radiation energy density becomes equal to the inflaton energy density, $\\rho_{\\phi}=\\rho_R$, and that standard reheating constraints on gravitino production can be applied. This may, however, overestimate the temperature at which the gravitino yield freezes out, as radiation does not yet fully dominate the energy density at this stage. Moreover, standard reheating constraints may not {\\it a priori} be applied in warm inflation scenarios due to the non-negligible abundance of gravitinos produced during inflation, which may potentially lead to a larger yield. This may, in fact, be the case also in conventional models, since the reheating phase is not necessarily instantaneous and thermal production of gravitinos may potentially occur for the duration of reheating and not freeze out until the universe is fully radiation dominated, resulting in a cumulative effect similar to that of warm inflation. Finally, we note that supersymmetry is broken during inflation, leading to gravitino masses parametrically close to the Hubble parameter and potentially to massive gauginos, which may also modify the production rate during warm inflation. Similarly, this may change the standard reheating constraints, as the gravitino mass also varies during the oscillating phase.   With these new insights in mind, we revisit the production of gravitinos in supersymmetric warm inflation, numerically evolving the Boltzmann equation for gravitinos into the radiation era. In Section \\ref{SGC} we give a brief review of the standard gravitino cosmology in cold inflation and in Section \\ref{WI} we outline the basic features of warm inflation, focusing as a working example on  monomial potentials in the sub-planckian regime. We discuss thermal gravitino production in warm inflation in Section \\ref{Gravprod} and present results for stable and unstable gravitinos, considering the effects of inflaton-dependent gaugino masses in both cases. In Section \\ref{Discussion} we summarize our main results and discuss possible directions of future research in this topic. ",
        "conclusions": "\\label{Discussion}  In this work, we have revisited the gravitino problem in warm inflation, focusing on thermal production which, providing the main difference from standard or cold inflation, places the strongest constraints on warm inflation dynamics. By performing a full numerical evolution of the gravitino yield into the radiation era we improve upon previous analyses. Firstly,  in the context of thermal gravitino production, the effective reheat temperature is the temperature at which the gravitino yield freezes out and not the temperature at which the inflaton energy density equals the radiation energy density. This allows the temperature to drop by approximately an order of magnitude, which lowers the final temperature at which gravitinos are produced compared to previous estimates. Secondly, we found that an analysis similar to standard reheating is in fact inadequate in describing gravitino production, due to the non-negligible yield produced throughout the whole duration of warm inflation. Finally, we have also taken into account the enhance particle masses during inflation due to supersymmetry breaking, in particular the gravitino and potentially the Standard Model gauginos.  Taking all of these issues into account, our work shows, in particular, that the final gravitino yield is substantially lowered for stronger dissipative effects, as in practice this lowers the temperature of the radiation bath during warm inflation significantly. We have presented regions of parameter space where the LSP gravitino can satisfy the dark matter bound and, for an NLSP gravitino, we determined the regions where the LSP abundance does not exceed the amount of dark matter present in our universe  and have given values of the dissipation parameter $C_{\\phi}$ for which late decays do not spoil the predictions of BBN.  Although thermal production is the dominant source of gravitinos during warm inflation, other non-thermal mechanisms may play a role at a later stage. Gravitinos can, in particular, also be produced from particle decays, but due to the large Hubble parameter during warm inflation these decays will not take place until the radiation era, at which point the standard cosmological results can be used. They can also be produced from the direct decay of the inflaton field, although we have found that, in the sub-planckian regime, the dissipative ratio $Q$ is necessarily large, which prevents the inflaton field from entering an oscillating phase. Figure \\ref{Oscillations} shows the inflaton field evolution as we artificially switch off dissipation at $\\rho_R=\\rho_{\\phi}$, at which point oscillations immediately begin.  \\begin{figure}[htbp] \\centering \\includegraphics[scale=1]{Oscillations.eps} \\caption{Switching off dissipation at $\\rho_R=\\rho_{\\phi}$, showing that the large dissipation keeps the inflaton field from oscillating in the radiation era. The dashed (solid) line corresponds to the case with (without) dissipation.} \\label{Oscillations} \\end{figure}  Dissipation will actually switch off when the heavy fields are no longer kinematically allowed to decay into the light degrees of freedom, which depends on their low scale mass hierarchy. For example, if supersymmetry is indeed a solution to the gauge hierarchy problem, we may expect light scalar masses to lie close to the TeV scale and dissipation to switch off at temperatures of this order. Apart from these kinematical constraints, the form of the dissipation coefficient in Eq.~(\\ref{Upsilon}) may actually hold down to very low temperatures. For example, to avoid exceeding the dark matter bound for the LSP gravitino, we require $C_{\\phi}\\sim1.5\\times10^{10}$ and for 40 e-folds of inflation, if the coupling $g\\sim 1$, the system remains in the low-temperature regime down to $T\\sim10\\,$MeV, at which point $\\rho_R/\\rho_{\\phi}\\sim10^{12}$. It is therefore unlikely in this case that any oscillations of the inflaton field may come to play a significant role in gravitino or, in fact, any entropy production. In particular, a significant dilution of the gravitino yield through a late inflaton decay along the lines proposed in \\cite{Bastero-gila} may be difficult to attain, although this may depend on the form of the inflaton potential, which goes beyond the scope of this work.  Our analysis revealed that it is possible to satisfy the dark matter constraint for LSP gravitinos and LSPs produced from NLSP gravitinos at large values of the dissipation parameter $C_{\\phi}$, which requires large couplings and field multiplicites, pointing towards beyond the Standard Model scenarios. The gravitino problem is more severe for unstable gravitinos potentially spoiling the predictions of BBN, and in this case much larger values of $C_\\phi$ are required.  One should note that such large values of the dissipation coefficient are nevertheless required in order to overcome the severe eta-problem affecting monomial potentials for sub-planckian values. Above the Planck scale, the potential gets exponentially steeper with increasing field values, requiring larger values of $C_{\\phi}$ to obtain 40-60 e-folds of inflation and also to suppress the resulting gravitino abundance. Using the full supergravity potential in Eq.~(\\ref{scalar_potential}) places a lower bound of $C_{\\phi}\\gtrsim10^8$ and $C_{\\phi}\\gtrsim2\\times10^7$ for 40 e-folds of inflation in the quartic and quadratic potentials, corresponding to $\\phi_*\\sim m_p$. Of course a non-canonical choice for the K\\\"ahler potential may alleviate this eta-problem, but supergravity is in any case unlikely to be the complete theory near the Planck scale and so any analysis along these lines must be taken with a pinch of salt. It should nevertheless be emphasized that simple monomial potentials cannot yield the required number of e-folds for sub-planckian values without dissipation, which is an attractive feature of warm inflation despite the large field multiplicities and/or couplings required.  One should bear in mind that observations may pose some constraints on the amount of dissipation present when the relevant CMB scales exit the horizon during inflation. In particular, an earlier analysis of non-gaussian effects on the primordial power spectrum showed that these depend logarithmically on the dissipative ratio at horizon crossing, $Q_*$, placing a model-dependent upper bound on the parameter $C_\\phi$ \\cite{Moss2007}. This analysis assumed, however, a constant dissipation coefficient, and more recently it was shown that for a generic $T$-dependence the non-gaussian parameter $f_{NL}$ is largely independent of the value of $Q_*$ in the strong dissipative regime, yielding $f_{NL}\\sim\\mathcal{O}(10)$ within the observable window of Planck \\cite{Moss:2011qc}. Hence, although the dynamics of second-order perturbations in warm inflation is not yet fully understood, we do not expect non-gaussianity to pose any significant constraints on our results.  In this work, we have considered a general scenario where all the MSSM degrees of freedom are in thermal equilibrium during inflation. However, it has been pointed out in \\cite{Rosa2012} that, in the low-temperature regime, $m_X\\ll T$, fermionic degrees of freedom may actually not thermalize, as both their contribution to the dissipation coefficient and their thermal scattering cross section are supressed compared to scalar fields. This is related to the structure of the superpotential (\\ref{twostage}) and the broken supersymmetry during inflation, which imply that the light fermions in the $Y$ multiplets only interact via the heavy $X$ bosons and fermions, whereas the light scalars have unsupressed interactions. Moreover, although the effects of gauge fields and their superpartners on the dissipation coefficient have yet to be analyzed in detail, their contributions to the dissipation coefficient may also be suppressed for sufficiently small gauge couplings. This would imply a thermal bath concurrent with inflation essentially composed of scalar particles, which would prevent gravitino production during inflation and eliminate the cumulative effect observed in our numerical simulations, at the same time requiring somewhat lower values of $C_\\phi$ for sub-planckian inflation. Both fermionic and gauge degrees of freedom will nevertheless be `reheaten' after inflation with either the exit from the low-temperature regime or the Hubble parameter dropping sufficiently in the radiation era. Although it requires further investigation, this may occur only at very low temperatures, as discussed above, in which case thermal gravitino production will be negligible.  In cold inflation there is a tension between having a large enough reheat temperature for thermal baryogenesis/leptogenesis to occur and it being low enough to avoid overproduction of gravitinos and other unwanted relics (see e.g. \\cite{Cline:2006ts}). In warm inflation this can be alleviated, as a baryon asymmetry can be produced through dissipation itself \\cite{Bastero-Gil2011}. Dissipation is an inherently out-of-equilibrium process, so the inclusion of baryon number and CP-violating interactions in the $X$ and $Y$ sectors in the superpotential Eq.~(\\ref{twostage}) naturally leads to the production of a baryon asymmetry during inflation. In the low-temperature regime, the produced asymmetry is naturally small despite the large couplings and field multiplicities required for a sufficiently long period of accelerated expansion and, moreover, this may lead to distinctive baryon isocurvature perturbations in the CMB anisotropies spectrum that may be observable in the near future. Warm inflation thus exhibits several attractive features that address not only the problems of inflationary dynamics itself but also many of the associated cosmological puzzles.  We would like to point out that the results from our analysis have a certain amount of crossover with cold inflation. Standard reheating is unlikely to be instantaneous and so the production of gravitinos will occur for the duration of the reheating phase. This will lead to an accumulated abundance similar to the one we have observed in warm inflation and so may change the standard reheating temperature constraints. The gravitino gets a mass from inflation and so, when the inflaton is oscillating about its minimum, the gravitino mass will also change at the same rate. If the oscillations are adiabatic, $\\dot{m}_{\\tilde{G}}/m_{\\tilde{G}} \\lesssim \\Gamma_{{\\text{scattering}}}$, then this effect can be analysed for various potentials in a similar way to the analysis performed in this work. It may then result in significant differences in the thermal production of gravitinos during the standard reheating picture.  With this work, we hope to have shed some light on gravitino production in warm inflation, with the way now paved for other potentials and dissipative coefficients to be analysed. In particular, the fact that inflation gives a mass to the gravitino may have a more significant impact on the thermal production in other potentials. Our analysis also brought to light some issues that may be significant to standard reheating and we hope that this motivates further exploration of this topic."
    },
    "1208/1208.4937_arXiv.txt": {
        "abstract": "Spherical Couette flow (flow between concentric rotating spheres) is one of flows under consideration for the laboratory magnetic dynamos. Recent experiments have shown that such flows may excite Coriolis restored inertial modes. The present work aims to better understand the properties of the observed modes and the nature of their excitation. Using numerical solutions describing forced inertial modes of a uniformly rotating fluid inside a spherical shell, we first identify the observed oscillations of the Couette flow with non-axisymmetric, retrograde, equatorially anti-symmetric inertial modes, confirming first attempts using a full sphere model. Although the model has no differential rotation, identification is possible because a large fraction of the fluid in a spherical Couette flow rotates rigidly. From the observed sequence of the excited modes appearing when the inner sphere is slowed down by step, we identify a critical Rossby number associated with a given mode and below which it is excited. The matching between this critical number and the one derived from the phase velocity of the numerically computed modes shows that these modes are excited by an instability likely driven by the critical layer that develops in the shear layer staying along the tangent cylinder of the inner sphere. ",
        "introduction": "Large-scale flows in stars or planets in many circumstances take place in a spherical shell. Most astrophysical fluid flows are also under the dominating influence of a background rotation. This rotation leads to the presence of inertial oscillations for which the restoring mechanism {comes from the} Coriolis acceleration.  The properties of these modes of oscillation are not fully understood, in part because they obey a hyperbolic equation in the space variables and therefore do not easily comply with boundary conditions. The solutions of this equation, known as the Poincar\\'e equation, have been studied in some details in the recent years, using numerical and analytical tools \\cite[e.g.][]{ML95,RV97,RGV01,RVG02,O05,RV10}. It has been shown that most of the eigenmodes of a rotating spherical fluid layer require viscosity to exist.  Indeed, viscosity is necessary to regularize the singularities formed by the focussing of characteristics by the boundaries. Viscosity transforms these singularities into shear layers whose thickness scales with some fractional power of it (exponents 1/4 or 1/3 are the most common). In addition, very recent works \\cite[][]{GL09,RV10} showed that the critical latitude singularity where the characteristics of the hyperbolic equation are tangent to the inner core boundary, plays a crucial role in periodically forced flows. The reason for that is not clear presently.  However, all the aforementioned previous works are theoretical studies considering idealized situations, and therefore should be compared to experimental studies.  Observations of inertial modes are not extensive, either in Nature or in the laboratory. A landmark in the experimental studies is the work of Aldridge\\cite{AlToo69,Ald72}. More recently, attractors of characteristics triggering oscillatory shear layers have been investigated experimentally for the understanding of ocean dynamics. Such experiments were conducted both on stably stratified fluids \\cite[][]{MBSL97,HBDM08,HTM10} and rotating fluids \\cite[][]{Maas01,MM04,Maas05} since internal modes (gravity waves) and inertial modes share many of the same mathematical properties. Other experiments have demonstrated the excitation of inertial waves in a fluid inside a precessing spheroidal cavity \\cite[][]{Noir2001a}.  \\begin{figure} \\centerline{ \\includegraphics[width=\\linewidth,clip=true]{fig/figure1.eps}} \\caption[]{Schematic of the 3-meter spherical-Couette apparatus. Inner and outer spheres rotate independently, driven by two 250 kW motors (not shown). There are three pressure sensors (azimuthally $90^\\circ$ apart) on the top lid ports, and an ultrasound velocimetry transducer measuring vertical components of fluid velocities (ultrasound beam depicted coming from port D).} \\label{expsetup} \\end{figure}  \\begin{figure} \\centerline{ \\includegraphics[width=\\linewidth,clip=true]{fig/figure2.eps}} \\caption[]{(Color online) Spectrogram from pressure measurements as the inner sphere rotates with different speeds in counter-rotation (inner sphere rotating in the opposite direction as the outer sphere). Each vertical line in the spectrogram is the power spectral density of the pressure using $\\rho \\RO \\Omega^2 L^2$ as the unit pressure. Outer sphere rotation rate is 1.5 Hz corresponding to $E=2.5\\times10^{-8}$. Modes indicated correspond to full sphere modes characterized by $(l,m,\\omega/2\\Omega)$.} \\label{spectrogram1} \\end{figure}  \\begin{figure} \\centerline{ \\includegraphics[width=\\linewidth,clip=true]{fig/figure3.eps}} \\caption[]{(Color online) Same as Fig.~\\ref{spectrogram1} but for a $\\RO$ range corresponding to co-rotation (inner sphere rotating in the same direction as the outer). The band near $\\omega/2\\Omega =0.05$, which has an azimuthal number $m=1$, is possibly not a single inertial mode. It's signature in the experiment reported by Kelley et al. \\cite[][]{KTZL10} is weak, perhaps a consequence of its quasi-geostrophic character (see discussion at end of Sect. IIIB).}  \\label{spectrogram2} \\end{figure}   All these experiments have shown that inertial modes are robust features of rotating fluid flows. In a very recent experiment aimed at studying a fluid dynamo, inertial modes were detected through their coupling with an imposed magnetic field, in a spherical Couette flow \\cite[][]{KTZL10}. This flow can indeed produce magnetic fields at sufficiently high magnetic Reynolds numbers \\cite[][]{GC10}. In that experiment, the fluid was contained in a spherical shell with inner and outer radii equal to 10~cm and 30~cm respectively. The Ekman number, the non dimensional measure of viscosity (see below), was approximately 10$^{-7}$.  In astrophysical or geophysical {(Earth's core)} situations this number is rather {less} than 10$^{-12}$. In fact, as was shown by recent numerical work \\cite[e.g.][]{RGV01,RV10}, the asymptotic regime describing vanishingly small viscosities usually appears at Ekman numbers below 10$^{-8}$. A new experiment geometrically similar to the Earth's core, using a sphere with an outer radius of 1.46m offers a unique opportunity to observe some near singular inertial modes close to their asymptotic regime since now Ekman numbers can be as low as $2.5\\times 10^{-8}$. Indeed, first results on this experiment using a  precessionally forced flow \\cite[][]{TZL12} provided a clear evidence of detached shear layers spawned by the critical latitude singularities.  \\begin{figure*} \\centerline{ \\includegraphics[width=0.50\\linewidth,clip=true]{fig/figure4a.eps} \\includegraphics[width=0.50\\linewidth,clip=true]{fig/figure4b.eps} } \\caption[]{(Color online) Rossby number dependence of the frequency of the two most prominent m=1-modes. The insert gives the equation of the linear best fit.} \\label{omega_ro1} \\end{figure*}  \\begin{figure*} \\centerline{ \\includegraphics[width=0.33\\linewidth,clip=true]{fig/figure5a.eps} \\includegraphics[width=0.33\\linewidth,clip=true]{fig/figure5b.eps} \\includegraphics[width=0.33\\linewidth,clip=true]{fig/figure5c.eps} } \\caption[]{(Color online) Same as Fig.~\\ref{omega_ro1} but for the three most prominent m=2-modes.} \\label{omega_ro2} \\end{figure*}  In this paper we further consider the results of this experiment. Thanks to a simple model based on numerical solutions of forced inertial modes in a spherical shell, we find a scenario that explains the excitation of inertial modes in a nonlinear spherical Couette flow.  For this, we first describe the experimental set-up and the observational facts concerning the observed inertial modes (sect.2). We then compute the response of an incompressible fluid inside a rigidly rotating spherical shell when some periodic forcing is applied (sect.3). We use this simple model to identify the resonance peaks and to interpret their full width at half-maximum (FWHM). We then discuss a scenario that explains most of the experimental facts (sect. 4). Conclusions end the paper. ",
        "conclusions": "In this work we have been able to refine the scenario proposed by Kelley et al. for the selection of inertial modes in a spherical Couette flow \\cite[][]{KTZL10}. We showed that the observed modes are most likely excited by a critical layer lying inside the shear layer which separates the fluid inside and outside the tangent cylinder. We demonstrated that this mechanism leads to a negative critical Rossby number, below which some non-axisymmetric retrograde inertial modes can be excited. The predicted value of this critical Rossby number matches quite nicely the observed experimental values, thus giving support to the proposed mechanism. However, because of the strong differential rotation needed to excite some modes, the identification of a few oscillation frequencies with those of a fluid inside a  spherical shell in solid body rotation remains uncertain. Another pending question is that of the equatorial symmetry of the observed modes. Our scenario along with the matching of frequencies of many modes argue in favour of a selection of anti-symmetric modes. However, this conclusion is not completely firm because some of the symmetric modes may match a few of the frequencies and also because the FWHM of two resonances show experimental values that are much less than those predicted by the model, which is {\\it a priori} less dissipative.  These results underline the need of more detailed investigations of inertial oscillations within differentially rotating fluids.  Preliminary numerical results \\cite[][]{BR12} show that a global shear can indeed significantly modify the properties of inertial modes and also shows that critical layers can destabilize some non-axisymmetric modes. However, the very case of the mean flow associated with a quasi-turbulent spherical Couette flow remains to be investigated.  Such a study may lead to the interesting perspective of reconstructing the interior differential rotation of the fluid by adjusting the prediction of the model to the observed value of the resonance frequencies. One could thus deduce details of the mean flow. Such a technique is similar to helioseismology techniques in solar research \\cite{thompson_etal96}, but instead using inertial modes to reconstruct the internal rotation of the fluid.  Inertial modes are certainly {the most appropriate} modes to infer the rotational property of a star or a planet. Our comparison of the inertial frequencies of our model with the observed ones, shows that a laboratory spherical Couette flow offers a unique playground to test this use of inertial waves.  Finally, let us point out that the ultimate challenge of modelling (this) experimental spherical Couette flows is to predict the amplitude of excited modes as a function of the Rossby number and to reproduce the sequence of their appearance."
    },
    "1208/1208.4748_arXiv.txt": {
        "abstract": "Primordial nucleosynthesis, or \\bbn\\ (BBN), is one of the three evidences for the Big-Bang model, together with the expansion of the Universe and the Cosmic Microwave Background. There is a good global agreement over a range of nine orders of magnitude between abundances of \\hli\\  deduced from observations, and calculated in primordial nucleosynthesis. This comparison was used to determine the baryonic density of the Universe. For this purpose, it is now superseded by the analysis of the Cosmic Microwave Background (CMB) radiation anisotropies. However, there remain, a yet unexplained,  discrepancy of a factor 3--5, between  the calculated and observed lithium primordial abundances, that has not been reduced, neither by recent nuclear physics experiments, nor by new observations. We review here the nuclear physics aspects of BBN for the production of \\hli, but also \\six, \\neu, \\onz\\ and up to CNO isotopes. These are, for instance, important for the initial composition of the matter at the origin of the first stars. Big-Bang nucleosynthesis, that has been used, to first constrain the baryonic density, and the number of neutrino families, remains, a valuable tool to probe the physics of the early Universe, like variation of \"constants\" or alternative theories of gravity. ",
        "introduction": "\\label{s:intro} \\bigskip   There are presently three evidences for the Big-Bang Model : the universal expansion, the Cosmic Microwave Background (CMB) radiation and Primordial or Big-Bang Nucleosynthesis (BBN). The third evidence for a hot Big-Bang comes, indeed, from the primordial abundances of the ``light elements'': \\hli. They were produced during the first $\\approx$20 minutes of the Universe when it was dense and hot enough for nuclear reactions to take place. These primordial abundances are compared to astronomical observations in primitive astrophysical sites. It is worth reminding that \\bbn\\ has been essential in the past, to first estimate the baryonic density of the Universe, $\\rho_{\\rm B} = (1-3)\\times10^{-31}$~g/cm$^3$ \\cite{Wag73}, and give an upper limit on the  number neutrino families $N_\\nu\\leq3$\\cite{Yan79}, both in the seventies. The number of light neutrino families is now known from the measurement of the $Z^0$ width by LEP experiments at CERN: $N_\\nu$ = 2.9840$\\pm$0.0082~\\cite{LEP}. The nuclear reaction rates have all been measured in nuclear physics laboratories or can be calculated from the standard theory of weak interactions. In that case, they are normalized to the experimental value for the lifetime of the neutron. Its precise value is still a matter of debate \\cite{Wie11} $\\tau_{\\mathrm{n}}$ = 880-884~s, but its uncertainty has only marginal effect on BBN. The last parameter to have been independently determined is the precise value of baryonic density of the Universe, which is now deduced from the observations of the anisotropies of the CMB radiation. It is usual to introduce $\\eta$, the number of photons per baryon which remains constant during the expansion, and is directly related to \\ob\\ by \\obh=3.65$\\times10^7\\eta$ with \\begin{equation} \\Omega_{\\mathrm{b}}{\\cdot}h^2=0.02249\\pm0.00062  \\; {\\rm and} \\; \\Omega_{\\mathrm{b}}=0.04455\\pm0.0027 \\label{eq:omega} \\end{equation} (``WMAP only Seven Year Mean'' \\cite{WMAP}). The parameter $h$ represents the Hubble constant ($H_0$) in units of 100~km/s/Mpc (i.e. $h$= 0.704  \\cite{WMAP}) and \\ob$\\equiv\\rho_{\\rm B}/\\rho_{\\rm 0,C}$ the baryonic density relative to the {\\em critical density}, $\\rho_{\\rm 0,C}$, which corresponds to a flat (i.e. Euclidean) space. It is given by : \\begin{equation} \\rho_{\\rm 0,C}={{3H_0^2}\\over{8{\\pi}G}}=1.88\\;h^2\\times10^{-29}\\; \\mathrm{g/cm^3}\\;\\mathrm{or}\\;2.9\\;h^2\\times10^{11}\\; \\mathrm{M_\\odot}/\\mathrm{Mpc^3} \\label{eq:critic} \\end{equation} where $G$ is the gravitational constant. It corresponds to a density of a few hydrogen atoms per cubic meter or one typical galaxy per cubic megaparsec (Mpc). This results (Eqs.~\\ref{eq:omega}--\\ref{eq:critic}) in a baryonic density which is just slightly above the range provided by Wagoner \\cite{Wag73} in 1973!   Hence, the number of free parameters in \\sbbn\\ has now been reduced to  zero, and the calculated primordial abundances are in principle only affected by the moderate uncertainties in some nuclear cross--sections. It may appears that \\bbn\\ studies are now useless, but this is certainly not the case. First, even though the agreement with observations is good or very good for \\qua, \\tro\\ and \\deu, there is a tantalizing discrepancy for \\sep\\ that has not yet found a consensual explanation. Second, when we look back in time, it is the ultimate process for which, {\\it a priori}, we know all the physics involved. Hence,  departure from its predictions could provide hints or constraints on new physics or astrophysics.   Besides the \\hli\\ isotopes, some minute traces of \\six, \\neu, \\onz\\ and CNO are produced by BBN. Observations of \\six\\ in a few halo stars have renewed the interest for this isotope and the nuclear uncertainties concerning its production. The evolution of the first generation (Population III) of stars could be influenced by the amount of primordial CNO elements, as hydrogen burning can proceed either through the slow p--p chain, or through the more efficient CNO cycle. We will hence review the nuclear aspects of the primordial production of element up to oxygen. ",
        "conclusions": "The baryonic density of the Universe as determined by the analysis of the CMB anisotropies is in very good agreement with Standard BBN compared to \\deu\\ primordial abundance deduced from cosmological cloud observations. However, it disagrees with lithium observations in halo stars by a factor that has increased with the availability of improved nuclear data and astronomical observations. Presently, the favored explanation is lithium stellar depletion, but the larger needed depletion factor is hardly compatible whit the thin observed plateau. It is hence essential to determine precisely the {\\em absolute} cross sections important for \\sep\\ nucleosymthesis (Table~\\ref{t:sensib}).  Nevertheless, primordial nucleosynthesis remains an invaluable tool for probing the physics of the early Universe. When we look back in time, it is the ultimate process for which we, {\\it a priori}, know all the physics involved. Hence, departure from its predictions provide hints for new physics or astrophysics. Gravity could differ from its general relativistic description, for instance a scalar field, in addition to the tensor field of general relativity (GR), appears naturally in superstring theories. That would affect the rate of expansion of the universe and hence BBN (see \\cite{Coc09} and Ref.~\\cite{JPU10} for a review). Coupled variation of the fundamental couplings is also motivated by superstring theories (see Ref.~\\cite{JPU11} for a review). However, the impact of these variations on the nuclear reaction rates is difficult to estimate, as in general, nuclear physics uses phenomenological models, whose parameters are not explicitly linked to fundamental constants. The decay of a massive particle during or after BBN could affect the light element abundances and potentially lower the \\sep\\ abundance (see e.g.~\\cite{Cyb10}). Negatively charged relic particle, like the supersymmetric partner of the tau lepton, could form bound states with nuclei, lowering the Coulomb barrier and hence leading to the catalysis of nuclear reactions (see e.g.~\\cite{Pos08,Kus08}). Annihilation of dark matter during BBN, e.g. injecting extra neutrons,  could also modify the primordial abundances \\cite{Jed04,Alb12}.  We have extended our network up to the CNO region and performed a sensitivity study to identify the few reactions that could affect the A$>$7 isotope yields and re-evaluated their rates. The CNO isotope production was found to be in the range CNO/H = $(0.5-3.)\\times10^{-15}$, not sufficient to have an impact on the evolution of the first stars. It is nevertheless a reference value for comparison with non-\\sbbn\\ CNO production e.g. in the context of varying constants. In this particular case, even with a faster triple--alpha reaction rate or a stable $^8$Be, the C(NO) production remains $\\approx$6 order of magnitude \\cite{Coc12a} lower than the \\sbbn\\ value reported here.   Last but not least, we stress here the importance of sensitivity studies in nuclear astrophysics that have been done, e.g. in the context of novae \\cite{Ili02}, X--ray burst \\cite{Par08} or massive stars \\cite{Ili11}. Even in the simpler context of BBN without the complexity (e.g. mixing) of stellar nucleosynthesis, it would have been very unlikely to predict the influence of  the $^1$H(n,$\\gamma)^2$H reaction on \\sep\\ nor of  the $^7$Li(d,n)2$^4$He reaction on CNO.    \\ack I am indebted to all my collaborators on these topics:  Pierre Descouvemont, Sylvia Ekstr\\\"om,  St\\'ephane Goriely, Georges Meynet, Keith Olive, Jean-Philippe Uzan, Elisabeth Vangioni and Yi Xu . This work  was supported in part by the french ANR VACOUL."
    },
    "1208/1208.4781_arXiv.txt": {
        "abstract": "{} {This paper is the first of a series investigating Lyman-alpha (hereafter \\lya) radiation transfer through hydrodynamical simulations  of galaxy formation. Its aim is to assess the impact of the interstellar medium (ISM) physics on \\lya\\ radiation transfer and to quantify how galaxy orientation with respect to the line of sight alters observational signatures. } { We compare the results of \\lya\\ radiation transfer calculations through the ISM of a couple of idealized galaxy simulations in a dark matter halo of $\\sim 10^{10}M_\\odot$. In the first one, G1, this ISM is modeled using physics typical of large scale cosmological hydrodynamics simulations of galaxy formation, where gas is prevented from radiatively cooling below 10$^4$K.  In the second one, G2, gas is allowed to radiate away more of its internal energy via metal lines and consequently fragments into dense star-forming clouds.  } {First, as expected, {\\it the small-scale structuration of the ISM plays a determinant role in shaping a galaxy's \\lya{} properties}. The artificially warm, and hence smooth, ISM of G1 yields an escape fraction of $\\sim50$\\% at the \\lya{} line center, and produces symmetrical double-peak profiles. On the contrary, in G2, most young stars are embedded in thick star-forming clouds, and the result is a~$\\sim10$ times lower escape fraction. G2 also displays a stronger outflowing velocity field, which favors the escape of red-shifted photons, resulting in an asymmetric \\lya{} line. Second, {\\it the \\lya{} properties of G2 strongly depend on the inclination at which it is observed}: From edge-on to face-on, the line goes from a double-peak profile with an equivalent width of $\\sim-5$\\AA{} to a 15 times more luminous red-shifted asymmetric line with EW $\\sim 90$\\AA{}.} {The remarkable discrepancy in the \\lya\\ properties we derived for two ISM models raises a fundamental issue. In effect, it demonstrates that \\lya\\ radiation transfer calculations can only lead to realistic properties in simulations where galaxies are resolved into giant molecular clouds. Such a stringent requirement translates into severe constraints both in terms of ISM physics modeling and numerical resolution, putting these calculations out of reach of current large scale cosmological simulations.  Finally, we find inclination effects to be much stronger for \\lya{} photons than for continuum radiation. This could potentially introduce severe biases in the selection function of narrow-band \\lya{} emitter surveys, and in their interpretation, and we predict these surveys could indeed miss a significant fraction of the high-$z$ galaxy population. } ",
        "introduction": "\\label{s_intro}  In the last decade, the Lyman-alpha (\\lya{}) emission line has become an observational tool of choice to detect high redshift galaxies via narrow-band surveys \\citep[e.g.][]{Hu1998,Kudritzki2000,Shimasaku2006,Ouchi2008,Ouchi2010,Hu2010} or blind spectroscopic searches \\citep[e.g.][]{vanBreukelen2005,Rauch2008,Cassata2011}. Today, the number of galaxies detected in this fashion (hereafter Lyman-Alpha Emitters, LAE) is becoming statistically significant, and LAEs play a major role in our census of high-$z$ galaxies. At the same time, spectroscopic follow-ups of UV-selected galaxies shed more and more light on the physical nature of LAE and on their place in the cosmic history of galaxy formation \\citep{Shapley03,Tapken07,Bielby11}. One of the major challenges in the years to come, both theoretical and observational, is yet to understand the details of the \\lya{} line profiles we observe: How do they relate (if they do) to any physical property of high-$z$ galaxies~?  Although a number of semi-analytic models for Lyman-alpha Emitting galaxies (LAEs) have been published \\citep[e.g.][]{Ledelliou06,Orsi08,Dayal2008,Dayal2009,Orsi2011,Dayal2011, Garel2012}, the complete radiation transfer through \\emph{the interstellar medium} of \\lya\\ emitting galaxies has been taken into account in only a handful of previous studies \\citep[][see Table \\ref{compar} for a summary.] {tasitsiomi06,Laursen2009, Barnes2012,Yajima2011,Yajima2012}.  \\begin{table*} \\begin{tabular}{c|c|c|c|c|c} & this study              & \\citet{tasitsiomi06}    & \\citet{Laursen2009}     & \\citet{Barnes2012}     & \\citet{Yajima2011,Yajima2012}       \\\\ \\hline context        & \\lya\\ emitting galaxies & \\lya\\ emitting galaxies & \\lya\\ emitting galaxies & DLA-host galaxies      & \\lya\\ emitting galaxies  \\\\ \\hline hydro technics & AMR (RAMSES)            & AMR (ART)               & SPH (TreeSPH)            & SPH (GADGET)           & SPH (GADGET)             \\\\ \\lya\\ RT       & \\lya\\ + continuum       & \\lya, no dust           & \\lya                    & \\lya, no dust & \\lya\\ + continuum        \\\\ &    AMR                  &   AMR                   & AMR                     & cartesian              & AMR                      \\\\ \\lya\\ sources  & recombination           & recombination           & recombination           & central point source   & recombination            \\\\ & from young stars        & from young stars        & + gravitational cooling &                        & + collisional excitation \\\\ &                         &                         & + UV background         &                        &                          \\\\ environment    & isolated galaxy         & cosmo zoom              & cosmo zoom              & cosmo zoom             & cosmo zoom               \\\\ nb of objects  & 2 & 1 & 9 & 3 & $\\sim 950$ \\\\ stellar mass   & $1.8\\times10^9$\\msun for G1 & $\\sim 10^{10}$\\msun & $6\\times10^6$ to         & $1.5\\times10^{10}$\\msun & $4.3\\times10^9$\\msun     \\\\ & $4.9\\times10^8$\\msun for G2 &                     & $3\\times10^{10}$\\msun    & $1.5\\times10^{11}$\\msun & $9.3\\times10^9$\\msun     \\\\ &                             &                     &                         & $7.5\\times10^{11}$\\msun & $4.1\\times10^{10}$\\msun   \\\\ stellar mass resolution & $1.4\\times10^3$\\msun for G2 & $2\\times10^4$\\msun & $10^6$\\msun           & not available          & $1.9\\times10^4$\\msun      \\\\ & $7.7\\times10^3$\\msun for G1 &                    &             &               & \\\\ spatial resolution$^{(a)}$& 18 pc for G2            & 29 pc               & 137 pc$^{(b)}$                  & 514 pc              & 342 pc                     \\\\ & 147 pc for G1           &                     &                         &                        &                            \\\\ \\hi\\ temperature   & $10^2$ to $10^5$K       & $10^3$ to $10^4$K    & $10^4$ K                & $10^{4.3}$ to $10^5$K   & -              \\\\ \\end{tabular} \\caption{Comparison of the 4 published studies of \\lya\\ radiation transfer through the interstellar medium of galaxies with this study. \\newline $^{(a)}$ Resolution, in physical $pc$. This is either the minimum cell size, for AMR codes, or the gas gravitational softening length for SPH codes. In both cases, this reflects the smallest scale onto which a gas overdensity may feel its own gravity. \\newline $^{(b)}$ This resolution corresponds to their S87 simulation, which best matches our halo mass.} \\label{compar} \\end{table*}  In \\citet{tasitsiomi06}, \\lya\\ radiation transfer is post-processed in the brightest \\lya\\ emitter of a gas-dynamics+N-body adaptive refinement tree \\citep[ART,][]{Kravtsov05} simulation at $ z\\sim 8$, in order to investigate its detectability. In \\citet{Laursen2009}, a sample of nine galaxies at $z = 3.6$ taken from cosmological N-body/hydrodynamical TreeSPH simulations \\citep{Sommer06} are considered, sampled in mass. The main result of the paper is an anti-correlation between \\lya\\ escape fraction and the mass of the galaxy. In \\citet{Barnes2012}, three halos at $z = 3$ are selected in cosmological hydrodynamic simulations (GADGET-2) aimed at reproducing the physical properties of the host galaxies of DLAs at $z\\sim 3$ \\citep{Tescari09}. In \\citet{Yajima2011,Yajima2012}, the same halo is followed at different redshifts, in order to study the evolution of the \\lya\\ properties of their galaxies with time.  These previous studies are done through a warm interstellar medium in which the gas is cooled down to T$\\sim10^4$K, with the consequence that the formation of small scale structures is not modelled. Indeed, the pressure support of this warm gas prevents it from collapsing at scales smaller than its Jeans length, explaining the difference in spatial resolution between the different experiments (see section \\ref{s_hydro} for more details). However, theoretical expectations suggest a strong dependance of the \\lya\\ transfer on the structure, and geometry of the interstellar medium of galaxies, and the main goal of our study is to investigate this point. Furthermore, two studies are dust-free, which prevent them from studying the \\lya\\ escape fraction from their configurations. Finally, monochromatic approaches of the problem do not allow to derive \\lya\\ equivalent widths, and compare the transfer of continuum versus line photons.  To overcome these limitations, we post-process hydrodynamical simulations of galaxy formation described in \\citet{Dubois08} performed  with the RAMSES code \\citep{Teyssier02}, with \\McLya{} \\citep{verhamme06}, including \\lya+continuum radiation transfer in a dusty medium, for two different ISM models : 1/ the reference model G1, comparable to previous studies, where the gas is cooled down to $10^4$K, 2/ a more realistic ISM model G2, where the gas is allowed to cool down to 100K, and the formation of small scale structures is followed.  The plan of this paper is the following. We start to describe the hydrodynamical simulations used to post-process \\lya\\ radiation transfer. Then, we describe the radiative transfer of \\lya\\ photons in the hydrodynamical simulations with \\McLya. The fourth part presents a comparison of the \\lya\\ properties for the two ISM models. In the fifth part, we discuss the effect of orientation on the \\lya\\ properties of G2, which appears as an interesting bonus result of this work. The sixth part describes the diffuse \\lya\\ halo around G2. The last part summarizes the main conclusions. ",
        "conclusions": "\\label{s_conclude}  In this paper, we have studied the \\lya{} properties of a couple of high-resolution simulated dwarf galaxies forming in an idealized dark matter halo.  Our two simulations assume different temperature floors of the cooling function (10$^4$K for G1, 100K for G2), which result in strikingly different structurations of the ISM. While the gas in G2 is able to fragment into small star-forming clumps, the thermal pressure support in G1 yields a rather smooth ISM with homogeneous star formation.  We have post-processed these galaxies with \\McLya{} in order to follow the resonant scattering of \\lya{} photons through their ISM, and to predict their resultant observational properties. Our main results are as follows. \\begin{itemize} \\item[1.] As expected, the small scale structure of the ISM plays a determinant role in shaping a galaxy's \\lya{} properties. In the G2 simulation, where gas is allowed to cool down to temperatures $\\ll 10^4$ K, most young stars are embedded in thick, dusty, star forming clouds, and the photons they emit are strongly attenuated. As opposed to the ``Neufeld scenario'', the {\\it clumpiness of the ISM here enhances the destruction of \\lya{} photons relative to continuum photons}. This is due to the fact that photons are emitted within the dense clouds, {\\it \\`a la} \\citet{CharlotFall00}, rather than outside, as assumed in \\citet{Neufeld91}.  In the G1 simulation, with an artificially warm ISM, young stars are found in much lower density environments and their photons escape more easily. This simulation is comparable to the previous studies in the literature which also include dust \\citep{Laursen2009,Yajima2011,Yajima2012}, and our results are indeed similar to these studies. Such simulations do not capture the enhancement of \\lya{} extinction relative to continuum in star-forming regions that we find in G2, simply because they do not form such dense star forming regions.  Another important feature is the kinematic structure of the ISM. Because G2 develops a genuine multiphase medium, with very dense star forming clouds and a tenuous diffuse component, supernovae explosions are able to push gas to high velocities (see Fig. \\ref{fig:velocity_field_g2}). Instead, the rather homogeneous ISM of G1 is overall denser than the diffuse medium of G2, and resists better supernovae explosions. It thus displays a rather static velocity field (see Fig. \\ref{fig:velocity_field_g1}). As shown in Sec. \\ref{s_spec}, these different velocity fields have a strong impact on the shape of \\lya{} lines.  \\item[2.] The analysis of \\lya{} emission from G2\\footnote{and G1, although we do not show it here} demonstrates the existence of {\\it a strong inclination effect}. Due to the numerous scatterings that line photons undergo, their probability to escape does not depend on the direction towards which they were emitted. Instead, they tend to systematically escape face-on, following the path of least opacity. Because continuum photons do not display such a strong angular redistribution, this effect is directly seen on the \\lya{} equivalent width, which we find to vary from $\\sim -5$ \\AA{} edge-on to $\\sim 90$ \\AA{} face-on.  We also find that this inclination effect is seen in the shape of the \\lya{} line emerging from our simulated galaxy. When seen edge-on, our galaxy has a double-peak line, associated with a low EW. When seen face-on, our galaxy has an enhanced red peak, and a high EW.  These results suggest the possible existence of strong observational biases in LAE surveys which necessarily rely on \\lya{} luminosity and EW selections, and could thus preferentially select face-on objects. As an example, a survey with an EW cut at $~20$~\\AA{} would select our galaxy only 40\\% of the times, assuming it has a random inclination.  \\item[3.] Scattering of galactic \\lya{} photons through the circum-galactic medium do produce an extended \\lya{} halo. We find that about a third of \\lya{} photons escaping G2 contribute to this diffuse component. This is somewhat at odds with the results of \\citet{Steidel2011}, though the comparison should be taken with caution given the fact that our simulated galaxy has a much smaller mass than the galaxies analyzed by these authors. Also, our simulation is idealized, and the CGM of G2 is not a good representation of what one finds in cosmological simulations at high redshift \\citep{RosdahlBlaizot2012, Dubois12}. \\end{itemize}  Although our quantitative results are clearly limited by much missing physics, we believe that this work demonstrates the two main following points: (1) resolving the ISM is mandatory if we want to understand the escape fraction of Lya from galaxies (it doesn\u2019t matter here if we have the correct solution: we show how widely the results vary when we change the ISM physics ... and this definitely shows that we should go further); and (2) for an ideal disc galaxy, we find that the escape fraction is a strong function of inclination, and we argue that this effect is quite possibly present in real galaxies (eventhough their morphologies are known to be more complex). Both results call for more work, both theoretically and observationally. From the theoretical viewpoint, we plan to make progress in forthcoming papers by (i) including the transfer of ionizing photons through a {\\it resolved} ISM, and (ii) embedding our galaxy in the full complexity of its cosmological context."
    },
    "1208/1208.6098_arXiv.txt": {
        "abstract": "We studied  the distributions of Si, Fe, and Ni in the intracluster medium (ICM) of the Coma cluster, one of the largest clusters in the nearby universe, using XMM-Newton data up to 0.5 $r_{180}$ and Suzaku data of the central region up to 0.16 $r_{180}$. Using the flux ratios of Ly$\\alpha$ line of H-like Si and 7.8 keV line blend to K$\\alpha$ line of He-like Fe, the abundance ratios of Si to Fe and   Ni to Fe   of the ICM were derived using APEC model v2.0.1. The Si/Fe ratio in the ICM of the Coma cluster shows no radial gradient. The emission weighted averages of the Si/Fe ratio in the ICM within 0.0--0.2 $r_{180}$, 0.2--0.5 $r_{180}$, and 0.0--0.5 $r_{180}$ are 0.97 $\\pm$ 0.11,  1.05 $\\pm$ 0.36 and 0.99 $\\pm$ 0.13, respectively, in solar units using the solar abundance table by \\citet{lodd2003}. These values are close to those of smaller clusters and groups of galaxies. Using the Suzaku data of the central region, the derived Ni/Fe ratio of the ICM is  0.6--1.5  in solar units, according to the same solar abundance table. The systematic difference in the derived abundance ratios by different plasma codes are about 10\\%. Therefore, for the ICM in the Coma cluster, the abundance pattern of Si, Fe, and Ni is consistent with the same mixture of the yields of supernova (SN) II and SN Ia in our Galaxy. Within 0.5 $r_{180}$, the cumulative iron-mass-to-light ratio increases with radius, and  its radial profile  is similar to those of relaxed smaller clusters with cD galaxies at their center. Using the observed Si/Fe ratio, the cumulative metal-mass-to-light ratios at 0.5 $r_{180}$ are compared with theoretical expectations. ",
        "introduction": "An important clue to the  evolution of galaxies is the metals in the intracluster medium (ICM). Because the Fe--K lines are prominent in the spectra of the ICM, the Fe abundance in the ICM has been studied in detail. With ASCA observations, \\citet{Fukazawa2000} found that clusters with a sharp X-ray emission centered on a cD galaxy commonly exhibit a central increment in the Fe abundance of the ICM. With Beppo-SAX observations, \\citet{deGrandi2001} also found a difference in the Fe abundance profiles between clusters with and without cool-cores. With XMM-Newton observations, \\citet{Johnson2011} found a similar trend among groups of galaxies. Within cool cores of clusters and groups of galaxies, the central Fe peak could have been mainly produced by type Ia supernovae (SNe) in cD galaxies. In contrast, during cluster merging, mixing of the ICM could destroy the cool cores and the central Fe peaks.  Since metals have been synthesized by SNe in galaxies, the ratios of  metal mass in the ICM to the total light from galaxies in clusters or groups, i.e., the metal-mass-to-light ratios, are the key parameters in investigating the chemical evolution of the ICM\\@. With ASCA observations, the derived  ratios of Fe mass in the ICM to the total light from galaxies, iron-mass-to-light ratio (IMLR),  within a radius where the ICM density falls below 3 $\\times 10^{-4} ~\\rm{cm^{-3}}$ is nearly constant in rich clusters and decreases toward poorer systems \\citep{Makishima2001}. In individual clusters, the IMLR is lower around the center \\citep{Makishima2001}. With Chandra, XMM, and Suzaku observations of groups and medium-size clusters, the lower IMLR  within a given over-density radius in some groups of galaxies have been confirmed \\citep{Matsushita2007a,  Tokoi2008, Rasmussen2009,  Sato2009a, Sato2009b, kSato2010, Komiyama2009, Sakuma2011,Murakami2011}.     Since Fe is both synthesized in SN Ia and SN II, to constrain contributions from the two types of SN, we need measurements of abundances of various elements. The ASCA satellite first studied the Si abundance in the ICM \\citep{Fukazawa1998, Fukazawa2000, Finoguenov2000, Finoguenov2001}. \\citet{Fukazawa1998} reported that the Si/Fe ratio in the ICM increases with ICM temperature, suggesting that the relative contribution of SN II increasing towards massive clusters. \\citet{Finoguenov2000} reported that the Si/Fe  ratio increases with radius in several clusters. Using Chandra data of groups out to $r_{500}$, \\citet{Rasmussen2007} found that SN II contribution increases with radius and completely dominates at  $r_{500}$. XMM-Newton and Suzaku observations also have been used to study the Si/Fe ratio of the ICM in clusters and  groups of galaxies. \\citep{Matsushita2003, Tamura2004, Sanders2006, Werner2006,  dePlaa2007, Rasmussen2007,Matsushita2007a, Matsushita2007,  Sato2007a, Sato2007b, Sato2008, Tokoi2008, Komiyama2009, deGrandi2009, Simionescu2009, Sato2009a, Sato2009b, kSato2010, Sakuma2011, Murakami2011}. With Suzaku observations of clusters and groups with the ICM temperatures lower than $\\sim4 $ keV,   the derived values of Si abundance are close to those of Fe  (the Si/Fe ratios are $\\sim$0.8 in solar units) out to 0.2--0.3 $r_{180}$, with a small scatter using the solar abundance table by \\citet{lodd2003}. In the core regions, the reported values of the Si/Fe ratios have a larger scatter \\citep{dePlaa2007, deGrandi2009}. With XMM-Newton observations of nearby clusters, \\citet{Tamura2004} found that the temperature dependence does not exist in the Si/Fe ratios. For example, the Si/Fe ratio outside the cool core of the Perseus cluster is 0.77 $\\pm$ 0.25  in solar units using the same solar abundance table. However, excluding cool core regions, the error bars in the Si/Fe ratio of hotter clusters are very large.      With Beppo-SAX observations, \\citet{deGrandi2002} reported that cooling flow clusters show higher  Ni abundances in their cores than non-cooling flow clusters. \\citet{B2005} analyzed  hundreds of clusters observed with ASCA and reported that the average Ni abundance in hot clusters is about 1.3 solar, which is significantly higher than the Fe abundances in the ICM. In the cool core of the Perseus cluster, the Ni/Fe ratio is consistent with the solar ratio \\citep{Churazov2004, tamura2009, Matsushita2011b}, whereas that of the Centaurus cluster is significantly higher \\citep{Matsushita2011b}. \\citet{dePlaa2007} found that the weighted average of the Ni/Fe ratios of core regions of nearby clusters was 1.4 $\\pm$ 0.3 with respect to solar ratio by \\citet{lodd2003}. \\citet{deGrandi2009} studied Si, Fe, and Ni abundances in the central regions of 26 local  clusters and discovered that the Ni/Fe ratio scatters significantly. In our Galaxy, [Ni/Fe] of stars are $\\sim$ 0, with no dependence of [Fe/H] \\citep{Edv1993, feltzing1998, Gratton2003}. This result indicates that in our Galaxy both SN II and SN Ia synthesize Ni in a similar manner as that of Fe. In cool cores, the enhancement of Fe abundance indicate the metal production from cD galaxies (e.g. \\cite{Hans2004}). Therefore, the higher Ni/Fe ratios observed in some clusters indicate that Ni synthesis in cD galaxies  differ from that in our Galaxy. However, in some cases, Ni abundances are derived from the spectral fitting including the Fe-L energy band. Then, Ni abundance is derived from residuals  between data and the Fe-L model, which strongly depends on the plasma codes. With CCD detectors, the K$\\alpha$ lines of He-like Ni and K$\\beta$ lines of He-like Fe are blended into a single bump at 7.8 keV, and  the derived Ni abundance  couples with  the effect of  resonant line scattering \\citep{Churazov2004, Matsushita2011b}. Therefore,  for cases in which  the optical depth for the scattering is sufficiently small, spectral fittings of K lines of Ni give more reliable Ni abundances.  The Coma cluster ($z=0.0231$), also known as Abell~1656, is one of the largest clusters in the nearby universe. The cluster does not have a strong cool-core in the center, and the X-ray peak is not associated with  two dominant galaxies in the central region \\citep{Vik94}. With Chandra observations,  \\citet{Vik01} found that these dominant galaxies retain their X-ray corona in the form of compact halo (a few kpc in size) with temperatures of 1--2 keV. Using XMM-Newton data, \\citet{arnaud01} found a temperature drop within 1 arcmin from one of the dominant galaxies, NGC 4874. However, they  reported that the projected temperature distribution of most of the core region is remarkably homogeneous. Using  Suzaku data,  \\citet{Sato2011} reported that fittings  of the continuum spectra and the ratio of the Ly$\\alpha$ line of H-like Fe and K$\\alpha$ line of He-like Fe   implied    the same ICM temperatures. Because line ratio is a steep function of temperature, this consistency supports the accuracy of temperature measurements using  Suzaku. \\citet{Matsushita2011} and \\citet{Sato2011} derived the Fe abundance profile in the ICM up to $\\sim$ 0.5 $r_{180}$. Within 0.2 $r_{180}$, the Fe abundance is flat at $\\sim$0.4 solar, according to the solar abundance table by \\citet{lodd2003}, and further decreases with radius. This flat Fe abundance   is significantly different from the peaked abundance profiles of cool-core clusters and indicates that the gases have been mixed well   in the core   during the past mergers associated with cluster growth.   In this paper, we study  Si/Fe and Ni/Fe ratios in the ICM of the Coma cluster observed with XMM and Suzaku. In addition, we derived  IMLR profiles of the Coma cluster and compared the results with smaller clusters. The paper is organized as follows; After the introduction, we present the observations in section 2, followed by the description of our data analysis and results in section 3. In section 4, we discuss our results.  We used the Hubble constant, $H_{\\rm 0} = 70$~km~s$^{-1}$~Mpc$^{-1}$\\@. The distance to the Coma cluster is $D_{\\rm L}=101$~Mpc, and $1'$ corresponds to 28~kpc. The virial radius of the Coma cluster, $r_{180}=1.95~h_{100}^{-1}\\sqrt{k \\langle T \\rangle/10~\\rm{keV}}~\\rm{Mpc}$ \\citep{Evrard1996, Markevitch1998}, is 2.5~Mpc for the average temperature $k \\langle T \\rangle$=7.8 keV. We used the solar abundance ratio by \\citet{lodd2003}, in which the solar Si, Ni and  Fe abundances relative to H are 3.47$\\times$10$^{-5}$,  1.66$\\times$10$^{-6}$, and 2.95$\\times$10$^{-5}$, respectively,  by number. Considering a difference in solar He abundance, the Fe abundance yielded by \\citet{lodd2003} is 1.5 times higher than that using the photospheric value by \\citet{angr}. Using the table by \\citet{lodd2003}, the Si/Fe and Ni/Fe ratios are factors of 1.55 and 1.48, respectively,  smaller than those from \\citet{angr}. Errors are quoted at 90\\% confidence level for a single parameter. The spectral analysis employed the XSPEC\\_v12.7.0 package. ",
        "conclusions": "We analyzed XMM (up to 0.5 $r_{180}$) and Suzaku (up to 0.16 $r_{180}$) data of the Coma cluster,  which is one of the largest clusters in the nearby Universe. Since the Coma cluster does not have   a  strong cool core   and the Fe abundance in the ICM is flat up to 0.2 $r_{180}$, the derived abundance pattern is not affected by recent metal supplies from the cD galaxies.  The Si/Fe ratios in the ICM are derived from the flux ratios of Ly$\\alpha$ line of H-like Si and K$\\alpha$ line of He-like Fe. The small temperature dependence of the line ratio limits the systematic uncertainty in the derived abundance ratio. The derived Si/Fe ratio in the ICM shows no radial gradient. The emission weighted average of the Si/Fe ratio within 0.0--0.2 $r_{180}$, 0.2--0.5 $r_{180}$, and 0.0--0.5 $r_{180}$ is 0.97 $\\pm$ 0.11, 1.05 $\\pm$ 0.36, and 0.99 $\\pm$ 0.13 in units of  solar ratio, using the solar abundance table by \\citet{lodd2003}. The comparison of the Si/Fe ratio of the Coma cluster with those of the smaller clusters indicates that dependence on the system mass of the Si/Fe ratio is small. The Ni/Fe ratio in the ICM is also derived from the flux ratio of the 7.8 keV line blend and K$\\alpha$ line of He-like Fe. The derived Ni/Fe ratio is 0.6--1.5 solar. Therefore,  the abundance pattern of Si, Fe, and Ni is consistent with the same mixture of the yields of SN II and SN Ia in our Galaxy.   We calculated the cumulative IMLR up to 0.5 $r_{180}$ using K-band and B-band luminosities of galaxies. Considering the observed Si/Fe ratio, at 0.5 $r_{180}$ the metal mass-to-light ratio in the ICM is consistent with expected value using a Salpeter IMF. However, if the IMLR continues to increase with a radius beyond 0.5 $r_{180}$, a flatter IMF is necessary. The IMLR of the Coma cluster is similar to those of poor clusters with temperatures of 2--4 keV. These clusters may have universal metal enrichment histories."
    },
    "1208/1208.5114_arXiv.txt": {
        "abstract": "The environments surrounding nine Wolf-Rayet stars were studied in molecular emission. Expanding shells were detected surrounding these WR stars (see left panels of Figure 1). The average masses and radii of the molecular cores surrounding these WR stars anti-correlate with the WR stellar wind velocities (middle panels of Figure 1), indicating the WR stars has great impact on their environments. The number density of Young Stellar Objects (YSOs) is enhanced in the molecular shells at $\\sim$5 arcmin from the central WR star (lower-right panel of Figure 1). Through detailed studies of the molecular shells and YSOs, we find strong evidences of triggered star formation in the fragmented molecular shells (\\cite[Liu et al.  2010]{liu_etal12}). ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.2028_arXiv.txt": {
        "abstract": "It has been argued recently that the galaxy peculiar velocity field provides evidence of excessive power on scales of $50\\hmpc$, which seems to be inconsistent with the standard $\\Lambda$CDM cosmological model. We discuss several assumptions and conventions used in studies of the large-scale bulk flow to check whether this claim is robust under a variety of conditions. Rather than using a composite catalogue we select samples from the SN, ENEAR, SFI++ and A1SN catalogues, and correct for Malmquist bias in each according to the \\textit{IRAS\\/} PSCz density field. We also use slightly different assumptions about the small-scale velocity dispersion and the parameterisation of the matter power spectrum when calculating the variance of the bulk flow. By combining the likelihood of individual catalogues using a Bayesian hyper-parameter method, we find that the joint likelihood of the amplitude parameter gives $\\sigma_8=0.65^{+0.47}_{-0.35}$ (68 per cent confidence region), which is entirely consistent with the $\\Lambda$CDM model. In addition, the bulk flow magnitude, $v \\sim 310 \\kms$, and direction, $(l,b)\\sim (280^{\\circ} \\pm 8^{\\circ}, 5.1^{\\circ} \\pm 6^{\\circ})$, found by each of the catalogues are all consistent with each other, and with the bulk flow results from most previous studies. Furthermore, the bulk flow velocities in different shells of the surveys constrain ($\\sigma_{8}$, $\\Omega_{\\rm{m}}$) to be ($1.01^{+0.26}_{-0.20}, 0.31^{+0.28}_{-0.14}$), for SFI++ and ($1.04^{+0.32}_{-0.24}, 0.28^{+0.30}_{-0.14}$) for ENEAR, which are consistent with {\\it WMAP\\/} 7-year best-fit values. We finally discuss the differences between our conclusions and those of the studies claiming the largest bulk flows. ",
        "introduction": "\\label{vel_intro} The cosmic bulk flow is the streaming motion of the galaxies surrounding our Milky Way system, due to the gravitational pull of cosmic structure on large scales. In the gravitational instability paradigm, for a galaxy at position $\\mathbf{r}$, the peculiar velocity of an individual galaxy at time $t$ is given by \\citep{Peebles93} \\begin{equation} \\mathbf{v}(\\mathbf{r,}t)=\\frac{\\Omega_{\\rm{m}}^{0.55}H_{0}}{4\\pi } \\int d^{3} \\mathbf{r^{\\prime }}\\delta_{m}(\\mathbf{r^{\\prime },}t) \\frac{\\mathbf{r}- \\mathbf{r^{\\prime }}} {|\\mathbf{r}-\\mathbf{r^{\\prime }}|^{3}}, \\label{pecu_def} \\end{equation}% where $\\delta_{\\rm m}(\\mathbf{r})=(\\rho (\\mathbf{r})-\\overline{\\rho })/% \\overline{\\rho }$ is the density contrast at position $\\mathbf{r}$, $\\Omega_{\\rm{m}}$ is the fractional matter density, and $H_{0}$ is the Hubble constant.  The bulk flow is normally considered as an average over a sufficiently large volume, with some window function $w(r,R)$, so that the above linear perturbation theory is applicable.  This average is defined as \\citep{Juszkiewicz90,Nusser11a} \\begin{equation} \\mathbf{V}_{\\rm bulk}(\\mathbf{r,}t) =\\frac{\\int d^{3}\\mathbf{r}^{\\prime }\\mathbf{v}(\\mathbf{r}^{\\prime} \\mathbf{,}t)w(\\left\\vert \\mathbf{r}^{\\prime }-\\mathbf{r}\\right\\vert ,R)} {\\int d^{3}\\mathbf{r}^{\\prime }w(\\left\\vert \\mathbf{r}^{\\prime } -\\mathbf{r}\\right\\vert ,R)},  \\label{bulk_def} \\end{equation} where $\\mathbf{v}(\\mathbf{r,}t)$ is the 3-D peculiar velocity field at time $t$, defined in Eq.~(\\ref{pecu_def}). A complete investigation of bulk flows of nearby galaxies should measure the individual velocities of galaxies all over the observed volume. However, realistic observational techniques, such as the Tully-Fisher relation, only allow us to probe the radial component of the peculiar velocities of galaxies. In addition, most of the current observations can only cover a patch of sky with limited depth, leading to large uncertainties when interpreting the results.   Of course, none of these considerations are new.  There is already a large literature on the study of the peculiar velocity field, with particularly intense activity in the early 1990s \\citep[see overviews in][]{Burstein,Courteau93,Latham91,CosmicVF,Strauss95,CosmicFlows}. Investigating the relationship between velocities and densities has great potential for constraining cosmological parameters, and testing theories of gravity on large scales.  However, it has long been realised that the construction of appropriate catalogues is difficult, and that systematic effects can easily overwhelm statistical noise.   In attempting to overcome these observational limitations, there have been significant recent efforts in the community to reconstruct bulk flow moments from the limited data available, and to test their consistency with the $\\Lambda$CDM cosmology. One of the important issues lies in determining the proper weighting for individual galaxy velocities in a catalogue in order to obtain streaming motions.  Some of the published studies, such as \\cite{Sarkar07} and \\cite{Abate09}, focus on a weighting scheme that produces the maximum likelihood estimate of the bulk flow \\citep[see also][]{Watkins09}, which can minimise the measurement noise. However, this weighting depends on the particular survey geometry and statistical properties, which leads to a large uncertainty when interpreting the constraints from combined data sets.   \\cite{Watkins09} proposed another method of estimating the bulk flow of galaxy peculiar velocities. They focused on the problem of how realistic surveys can be used to reconstruct the bulk flow at a given depth. They developed a minimum variance weighting method \\citep{Watkins09,Feldman10}, which minimises the variance between the real data catalogue and the ideal survey, and they applied it to combined catalogues of peculiar velocity surveys. Surprisingly, they found a very large bulk flow on $50\\hmpc$ scales ($v=407\\pm 81 \\kms$) towards $l=287^{\\circ}\\pm 9^{\\circ}$, $b=8^{\\circ}\\pm 6^{\\circ}$, which prefers a large amplitude of fluctuations ($\\sigma_{8}$), inconsistent with the {\\it WMAP\\/} 5-year results \\citep{Komatsu09}. Subsequent work has discussed a possible explanation for this large bulk flow related to pre-inflationary isocurvature perturbations \\citep{Ma11}.   Contradicting the claim in \\cite{Watkins09}, \\cite{Nusser11a} developed a method termed the ASCE (All Space Constrained Estimate) which reconstructs the bulk flow from an all-space 3-D velocity field to match the inverse Tully-Fisher relation. By applying this method, as well as the Maximum likelihood method \\citep{Abate09}, to the Spiral Field $I$-band Survey (SFI++ survey, \\citealt{Springob07}) catalogue, \\cite{Nusser11a} found the bulk flow on a sphere of $40\\hmpc$ radius to be $v=333 \\pm 38 \\kms$, towards ($l,b$)=($276^{\\circ}\\pm 3^{\\circ} ,14^{\\circ}\\pm 3^{\\circ}$), which is close to the results from the maximum likelihood method. The estimated cosmological parameters, i.e.\\ $(\\Omega_{\\rm{m}},\\sigma_8)=(0.236,0.88)$, are consistent with the $\\Lambda$CDM model. However, since \\cite{Nusser11a} only used the SFI++ data set, it is still not clear whether it is the other data sets used in \\cite{Watkins09} which led to the significantly different results.   Any analysis which claims to strongly rule out the simple inflationary $\\Lambda$CDM model deserves careful scrutiny, since a confirmed discordance would have profound consequences for our understanding of the large-scale Universe.  We can identify four potential problems in \\cite{Watkins09} which may potentially skew the likelihood and bias the results. Firstly, the inhomogeneous Malmquist bias is not corrected for in most catalogues, for example: ENEAR \\citep{Costa00,Bernardi02,Wenger03}; SN \\citep{Tonry03}; SC \\citep{Giovanelli98}; EFAR \\citep{Colless01}; and Willick \\citep{Willick99}. This deficiency can significantly bias the distance estimates.  Secondly, the distance errors from the Tully-Fisher and Fundamental Plane methods can be comparable to the measured velocities as the surveys go deeper, and moreover a simple model of Gaussian errors is almost certainly inappropriate as systematics come to dominate the distance estimation. Therefore the velocity data beyond $100\\hmpc$ become both very noisy and unreliable in assessing the bulk flow. Thirdly, directly combining various catalogues with different calibration methods can also induce systematic errors and a spurious flow. Finally, the assumption of a unique small scale velocity dispersion $\\sigma_{\\ast}=150 \\kms$ may be too small for some of the surveys (e.g. SFI++ prefers $400 \\kms$, \\citealt{Ma11}), perhaps skewing the constraints on the cosmological parameters $\\sigma_{8}$ and $\\Omega_{\\rm{m}}$. The purpose of this paper is to investigate carefully the analysis presented in \\cite{Watkins09}, and to combine each catalogue with a Bayesian hyper-parameter method to test for consistency with the usual $\\Lambda$CDM perturbation theory. As we have seen from \\cite{Nusser11a}, different statistical methods should not dramatically alter the results, so we will focus on the `minimal variance' scheme \\citep{Watkins09,Feldman10}.   A further motivation for this paper is as an extension to the velocity-gravity comparison work we have already carried out in \\cite{Ma12}. In that paper we compare the observational peculiar velocity data with the reconstructed velocity field from the {\\it IRAS\\/} PSCz catalogue, and fit the linear growth rate parameter, $\\beta$. In this new paper, we do not discuss the small-scale modes, but will reconstruct the bulk motion of galaxies on distances ${\\sim}\\,50\\,h^{-1}$Mpc. We will perform a direct comparison between observational data and $\\Lambda$CDM model predictions for the bulk flow velocity, and constrain the cosmological parameters $\\sigma_8$ and $\\Omega_{\\rm m}$. In addition, we will extend the minimal variance scheme suggested in \\cite{Watkins09} and \\cite{Feldman10} to a multishell likelihood method. Furthermore, we will directly investigate the reason for the apparently large flows found in \\cite{Watkins09} and \\cite{Feldman10}.   This paper is organised as follows. We first list the data sets used in Section~\\ref{vel_data}, and then discuss the data selection criterion in Section~\\ref{dataselect}. For the selected data, we correct the inhomogeneous Malmquist bias for the distance estimate (Section~\\ref{mbcorrect}). In Section~\\ref{mv_scheme}, we first illustrate how to quantify the variance of the bulk flow at any particular depth (Section~\\ref{mean_s_v}), then we review the minimum variance weighting scheme proposed in \\cite{Watkins09} to measure the bulk flow at a given depth (Section~\\ref{vel_mv}), and furthermore we present the likelihood function for each individual catalogue (Section~\\ref{individual_like}) and the hyper-parameter approach used to combine different data sets (Section~\\ref{vel_bayes_combine}). Then in Section~\\ref{result_discuss} we compare our findings with those in \\cite{Watkins09}.  We first confirm that we can accurately reproduce the results in \\cite{Watkins09} by adopting the same conventions; then in Section~\\ref{bfmoment} we show our constraints on bulk flow moments by performing the full likelihood analysis for each individual catalogue rather than the combined catalogue. In Section~\\ref{cosmologypara}, we apply the Bayesian hyper-parameter method to combine the likelihoods of different catalogues, in order to avoid the systematics that may affect the constraints. This allows us to assess the consistency of each individual catalogue, and to work out the cosmological parameters in the combined likelihood. In Section~\\ref{correlation-depth}, we extend our likelihood analysis to consider bulk flows in multiple shells in a survey, and their covariance matrix, and we compare our findings with {\\it WMAP\\/} 7-year best-fit values and the results from \\cite{Nusser11a}. Our discussion and conclusion are summarised in Section~\\ref{vel_conclude}.   Note that although $H_0$ ($=100h\\,{\\rm km}\\,{\\rm s}^{-1}{\\rm Mpc}^{-1}$) is now determined with reasonable accuracy, throughout this paper we continue to adopt the convention of giving distances in units of $\\hmpc$ for ease of comparison with previous results. ",
        "conclusions": "\\label{vel_conclude}  \\begin{table*} \\begin{centering} \\begin{tabular}{@{}lll}\\hline & \\cite{Watkins09} & This study \\\\ \\noalign{\\vspace{3pt}} \\noalign{\\hrule} \\noalign{\\vspace{3pt}} Cosmological parameters & {\\it WMAP\\/} 5-year \\citep{Komatsu09} & {\\it WMAP\\/} 7-year \\citep{Komatsu11} \\\\ Small scale velocity dispersion $\\sigma_{\\ast}$ & $150 \\kms$ & $400 \\kms$ \\\\ $P(k)$ calculation & Parameterisation with $\\Gamma=\\Omega_{\\rm{m}}h$ & Numerical result from CAMB \\\\ Distance indicator & Malmquist bias uncorrected & Malmquist bias corrected \\\\ Catalogues & COMPOSITE: combination of & SN, SFI++ and ENEAR catalogues \\\\ & SBF, SN, SFI++, ENEAR& \\\\ & SC, SMAC, EFAR and Willick  & \\\\ Data selection & None & Trim to $d \\leq 80\\hmpc$, $|v| \\leq 3000 \\kms$ \\\\ Number of samples & COMPOSITE (4536) & SN (78), ENEAR (669), SFI++ (2404) \\\\ Combination method & Direct combination & Hyper-parameter likelihood $\\&$ Multi-shell likelihood \\\\ \\hline \\textbf{Result for normalisation} & $\\sigma_{8}=1.7 \\pm 0.28$  & $\\sigma_{8}=0.65^{+0.47}_{-0.35}$ (hyper-parameter) \\\\ & & $\\sigma_{8}=1.01^{+0.26}_{-0.20}$ (SFI++) \\text{ } $\\sigma_{8}=1.04^{+0.32}_{-0.24}$ (ENEAR) \\\\ & (excluded by {\\it WMAP\\/} at $99\\%$) & (consistent with {\\it WMAP\\/}) \\\\ \\hline \\end{tabular}% \\caption{Comparison of the methodology, data selection, and results of our constraints with those in \\citet{Watkins09}.} \\label{tab3} \\end{centering} \\end{table*}  In this paper, we have been investigating bulk flow measurements using various catalogues. We find results which are different to those given by \\cite{Watkins09}, who claimed evidence for a surprisingly large bulk flow on $50\\hmpc$ scales, apparently discrepant with the $\\Lambda$CDM prediction. In contrast, by carefully considering four selected catalogues, we find a coherent flow of about $300 \\kms$ on a scale of $50\\hmpc$, entirely consistent with the value expected given the {\\it WMAP\\/} 7-year cosmological parameters.   By employing the same weighting scheme and the same conventions, we are able to accurately reproduce the results in \\cite{Watkins09}, as shown in Fig.~\\ref{reproduce1}.  Since we focus on the SN, SFI++ and ENEAR catalogues, we removed the other sub-catalogues from the COMPOSITE catalogue, and found a slightly lower value of $\\sigma_8$ (red line in Fig.~\\ref{reproduce1}b), but still higher than the {\\it WMAP\\/} constraint. This indicates that the high value of $\\sigma_8$ inferred from the COMPOSITE catalogue is not completely driven by the five deep and sparse catalogues included (SMAC, SBF, SC, EFAR and Willick).  To summarise the various other issues which could be responsible for the discrepancy, in Table~\\ref{tab3} we list several technical points which lead to quantitatively different results.   The first issue is the assumption of small-scale velocity dispersion, which goes into the calculation of the covariance matrix (Eq.~(\\ref{noise_mat1})). \\cite{Watkins09} assumed a value of $150 \\kms$, which is too small compared to the constraint obtained by \\cite{Ma11,Ma11b}, which was closer to the $400\\kms$ we chose here.  Besides this, \\cite{Watkins09} used an inaccurate approximation for the matter power spectrum. From Fig.~\\ref{pkcompare1}, one can see that although this is a small effect, it has the same sign, yielding smaller flows. Thus, by fitting to the observed flows, this tends to further increase the normalisation parameter $\\sigma_{8}$.   The second major difference lies in the inhomogeneous Malmquist bias correction. In \\cite{Watkins09}, only the SFI++ and SMAC catalogues were corrected for this effect.  In our approach, we used the full-sky density field from the PSCz catalogue to extrapolate the density $n(r)$ at any spatial position, and calculate the probability of the true distance $r$ given the measured distance $d$ (Eq.~(\\ref{MBP1})).  The comparison between the measured distance/velocity and true distance/velocity in Fig.~\\ref{MBcorrect}, shows that the bias tends to move galaxies to smaller distances.   Another difference is that we only keep the high quality samples SN, SFI++ and ENEAR from the \\cite{Watkins09} compilation, and we further include the recent compilation of supernovae data, i.e.\\ the A1SN catalogue. To remove any possible bias from the distant and sparsely sampled region, we restricted our attention to $d \\leq 80\\hmpc$, and to avoid the results being driven by outliers, we also limited our samples to $|v|\\leq 3000 \\kms$.  Furthermore, rather than using the COMPOSITE catalogue, we combined individual sample likelihoods using the Bayesian hyper-parameter technique. This should avoid the possibility that inconsistent data sets may bias the result if they are assigned equal weight. From the hyper-parameter likelihood, we find the best-fit value $\\sigma_8=0.65^{+0.47}_{-0.35}$. This is somewhat low and hence inconsistent with a large bulk flow. However, the uncertainty is so large that this result is still consistent with standard $\\Lambda$CDM expectations.   Finally, we proposed a multishell likelihood method, which calculates the bulk flows in all shells within a certain radius together with their covariance matrix. This multishell likelihood takes into account the scale-dependence of the matter power spectrum $P(k)$ on the $\\Omega_{\\rm{m}}$ parameter, and therefore maximises the constraining power one can obtain from a data set. By applying this likelihood to the SFI++ and ENEAR catalogues, we showed that they can provide much stronger constraints on $\\Omega_{\\rm{m}}$ and $\\sigma_{8}$ than the single shell ($50 \\hmpc$) constraint. Our result also shows consistency with {\\it WMAP} 7-year best-fits and results from \\cite{Nusser11a}.   We conclude that the apparently large bulk flow on $50\\hmpc$ scales found by \\cite{Watkins09} may not be a genuine flow. By correcting for Malmquist bias, carefully selecting samples and examining assumptions, one finds that the current peculiar velocity field catalogues are consistent with the $\\Lambda$CDM model. On the other hand, any claimed discrepancy is not due to the `minimal variance' scheme proposed by \\cite{Watkins09} and \\cite{Feldman10}, since in our tests, we have shown that this scheme gives consistent results. In addition, our conclusions also agree with several other independent searches for bulk flows, such as the ASCE method with the SFI++ catalogue \\citep{Nusser11a}, the minimal variance method with the Type-Ia SN data \\citep{Turnbull12}, and the luminosity function method with the 2MRS samples \\citep{Branchini12}. It should also be pointed out that the lack of evidence for a bulk flow on $50 \\hmpc$ removes some of the support for an excessive flow $\\sim1000 \\kms$ on even deeper scales $\\sim300 \\hmpc$ \\citep{Kashlinsky08}.   It seems clear that, despite extensive effort for decades, peculiar velocity catalogues remain systematics dominated. By applying different, but apparently reasonable, assumptions and statistical approaches, it is possible to find quite discrepant results using essentially the same data sets.  This means that the realistic error bars are probably larger than given in many of the published studies. In addition, one should notice that there are many other methods developed to compute bulk flows that do not rely on distance indicators, such as luminosity fluctuations and fluctuations in the galaxy number density \\citep{Branchini12}, as well as the use of the kinetic Sunyaev-Zeldovich effect (e.g. \\citealt{Osborne11}). Although they also suffer from systematic effects, these will be of a different nature and therefore such approaches can be regarded as complementary to the method discussed here. Large-scale bulk flows still offer promise for constraining cosmological models, but fully realising that promise will require further improvements in the construction of catalogues, and in the control of the systematic effects which continue to plague this field.   \\vskip 0.1 truein   \\noindent \\textbf{Acknowledgments:} We would like to thank Michael Hudson and Stephen Turnbull for sharing with us the First Amendment Supernovae compilation, and Enzo Branchini and George Efstathiou for helpful discussions. This research was supported by the Natural Sciences and Engineering Research Council of Canada. YZM is supported by a CITA National Fellowship.  \\appendix"
    },
    "1208/1208.2819_arXiv.txt": {
        "abstract": "{ More than one hundred GeV pulsars have been detected up to now by the LAT telescope on the {\\em Fermi} gamma-ray observatory, showing peak energies around a few GeV. Current modelling proposes that the high energy emission comes from outer magnetospheric gaps, however radiation from the equatorial current sheet which separates the two magnetic hemispheres outside the light cylinder has also been investigated.} {In this paper we discuss the region right outside the light cylinder, or \"near wind\" zone. We investigate the possibility that synchrotron radiation emitted by thermal populations in the equatorial current sheet of the pulsar wind in this region can explain the lightcurves and spectra observed by {\\em Fermi}/LAT.} { We use analytical estimates as well as detailed numerical computation to calculate the $\\gamma$-ray luminosities, lightcurves and spectra of $\\gamma$-ray pulsars. } { Many of the characteristics of the $\\gamma$-ray pulsars observed by {\\em Fermi}/LAT can be reproduced by our model, most notably the position of these objects in the $P-\\dot{P}$ diagram, and the range of $\\gamma$-ray luminosities. A testable result is a sub-exponential cutoff with an index $b=0.35$. We also predict the existence of a population of pulsars with cutoff energies in the MeV range. These have systematically lower spindown luminosities than the {\\em Fermi}/LAT detected pulsars. } { It is possible for relativistic populations of electrons and positrons in the current sheet of a pulsar's wind right outside the light cylinder to emit synchrotron radiation that peaks in the sub-GeV to GeV regime, with $\\gamma$-ray efficiencies similar to those observed for the {\\em Fermi}/LAT pulsars.} ",
        "introduction": "Since the launch of the {\\em Fermi} space telescope, the sample of gamma ray pulsars has grown to include more than one hundred objects. All these pulsars, whether young or millisecond, single or in binaries, show spectra consistent with power laws with exponential cutoffs, where the cutoff energy lies in the range 1-10 GeV \\citep{latfirstyear}. A prominent exception is the Crab pulsar, which has been detected by ground-based Cerenkov arrays in the TeV regime, and the spectrum of which can be well fitted by a broken power law \\citep{crabmagic,mccannetal11}. These observations point to the outer gap/slot gap models as the most probable explanation for the pulsed gamma-ray emission, based on both population prediction statistics and lightcurve modelling \\citep{gonthieretal10,decesaretal11,pierbattistaetal11,venteretal11,watters+romani11,venteretal12}, although for millisecond pulsars one has to evoke non-dipolar field geometries or displaced polar caps, in order to reach the required energies \\citep{harding+muslimov11}. In those models, the emission originates within the light cylinder, defined by the cylindrical radius $\\rlc = c/\\omega$ from the pulsar's rotational axis, where a corotating particle would reach the speed of light. However it was pointed out by \\cite{bai+spitkovsky10} that the emission region must be extended slightly outside the light cylinder in order to reproduce the doubly-peaked light curves using the magnetic field configuration from self-consistent, force-free simulations of pulsar magnetospheres.  The idea that high energy pulsations might come from current sheets near or outside the light cylinder is not new. \\cite{lyubarsky96} predicted that particles accelerated though reconnection close to the light cylinder might emit gamma-rays through the synchrotron process. Lyubarsky's emission site is the point where the warped equatorial current sheet which separates regions of opposite magnetic field polarity in the pulsar's wind meets the current flowing to/from the pulsar's polar caps, the so-called Y-point \\citep{spitkovsky06}. Later it was pointed out that pulsed emission naturally arises from the periodicity of the pulsar wind, which is modulated by the star's rotation, in combination with the large bulk Lorentz factors of the outflow \\citep{kirketal02}. This model has been successfully applied to calculate the polarization of optical emission from the Crab pulsar \\citep{petri+kirk05} and it has also shown promise in explaining the gamma-ray light curves observed by {\\em Fermi}/LAT \\citep{petri11}. In these models the emission is attributed to the inverse Compton process and starts in the wind region far from the light cylinder, $r \\gg \\rlc$. In the same context it was proposed that the gamma-ray radiation of the {\\em Fermi}/LAT band could be synchrotron radiation from power-law electrons in the pulsar wind's equatorial sheet far from the light cylinder \\citep{kirketal02,petri12}.  There is, however, no obvious reason why emission should be truncated outside the light cylinder (for magnetospheric gap models) or start at a minimum radius in the far wind zone (for the striped wind model). The emission region should evolve continuously from the outer magnetospheric gaps through the light cylinder and into the  equatorial current sheet, possibly affected by the reconnection process that will inevitably occur in the sheet \\citep{bai+spitkovsky10}. The region beyond the light cylinder but still in the \"near wind\" zone is the area that we are trying to explore. The purpose of this paper is to demonstrate how emission in the 100~MeV-100~GeV range explored by {\\em Fermi}/LAT can naturally arise from thermal populations of particles in the striped wind, without going into the details of reconnection physics, but just using a few simple assumptions about the local description of the current sheet. In section 2 we describe the model we use, including our assumptions, and give some analytical estimates for the emitted radiation. In section 3 we present examples of emission maps and phase-averaged spectra which arise from our model and explore the parameter space in order to come to more general conclusions about the pulsar population observed by {\\em Fermi}/LAT. Finally we discuss our results and the potential of a more detailed description of the current sheet outside the light cylinder as a source of pulsed energetic radiation. ",
        "conclusions": "In this paper we have shown how gamma-ray pulses can naturally arise within the framework of the pulsar wind's equatorial current sheet outside (but close to) the light cylinder. The advantage of this emission model is that it is an intrinsic mechanism that naturally produces peak energies in the MeV-GeV range. It can give meaningful results for almost all pulsars, irrelevant of age or environment, employing only a few parameters: the bulk Lorentz factor of the outflow $\\Gamma$, the sheet thickness $\\Delta$,  the obliquity $\\chi$ and the angle of the rotational axis to the line of sight $\\zeta$. To these the minimum radius $R_{\\rm min}$ can be added (which, in the calculated examples, has been set equal to $R_{\\rm min}=1$). However it is important to note that $R_{\\rm min}$ has to be sufficiently close to the light cylinder, so that the peak energy of the spectrum reaches the GeV regime. The parameters of the model are constrained by the characteristics of the pulsar, and by the physics of the current sheet. The novelty of our model is that it makes an attempt to account for the previously ignored region of the wind between $R=1$ and $R\\simeq \\Gwind$, and is therefore able to take advantage of the strong poloidal field close to the light cylinder, something that has not been discussed in previous wind models.  The predictions of our model are the following: \\begin{enumerate} \\item Two pulses per pulsar period are expected. The pulses have the same amplitude. \\item The shape of the pulses is symmetric with respect to their peak, as long as emission starts close to the light cylinder, and become increasingly asymmetric with rising $R_{\\rm min}$. \\item There is no significant interpulse emission in the GeV range. \\item The width of the pulse decreases with increasing photon energy. \\item The separation of the peaks varies with the obliquity $\\chi$ and the angle to the line of sight $\\zeta$, as in previous models of pulsed radiation from the pulsar wind. \\end{enumerate}  The features of the pulsed emission that cannot be explained by our model, such as the interpulse or the peaks of different intensity, could be accomodated by a model in which radiation comes both from within the light cylinder, possibly in an outer gap, as well as from the beginning of the wind. The difficulty of outer gap models to produce double peaked lightcurves, as noted in \\cite{bai+spitkovsky10}, in combination with the difficulty of our near wind model to produce single pulses, point to the need for a combined model in which the high energy emission starts within the light cylinder and continues in the equatorial current sheet outside the light cylinder. Therefore the physics of the region of formation of the current sheet is very important in the endeavour to understand $\\gamma$-ray emission from pulsars.  For pulsars in binaries, the possibility exists that the favoured radiation process is not synchrotron but rather inverse Compton on the low frequency photons of the massive star, as investigated in \\cite{petri+dubus11}. In this case the peak energies and emitted luminosities by the two mechanisms can be comparable, depending on whether the comoving energy density of the low frequency photons in the wind frame is comparable to the magnetic field energy density close to the light cylinder. For the pulsar B1259-63, which is a member of a binary with a B2e star, synchrotron radiation close to the light cylinder would give a peak energy $\\mathcal{E}_{\\rm GeV}\\sim 3$ (assuming $\\Gamma=10$ and $\\Delta =0.01$), while inverse Compton radiation on the star's photon gives a peak energy that is constrained by the Klein-Nishina limit to $\\mathcal{E}_{\\rm GeV}\\sim 10$. In this case the low frequency photon energy density is several orders of magnitude lower than the one of the magnetic field, therefore synchrotron radiation should be the dominant component. This, however, is not necessarily the case for other binary systems, such as LS 5039, where the seed photon energy density exceeds the one in PSR B1259-63 by a factor of 200. The detection of orbital modulation of the received $\\gamma$-ray flux can be the criterion on which to distinguish between the two emission mechanisms.  A possible issue with our model is that the supersonic solution of \\cite{bogovalov99} might not apply for the objects of lower bulk Lorentz factors. Particularly, the wind can accelerate, as is predicted in various magnetohydrodynamic models of pulsar winds \\citep{beskinetal98,kirketal09}. In this case the shape of the current sheet can in reality be different, since there is generally also a polar magnetic field component, which in our model was zero. However, as was mentioned above, the very steep dependence of $\\mathcal{E}_{\\rm GeV}$ on $R$ implies that even if the parameters $\\Delta$ and $\\Gamma$ changed with $R$, the present model would still be able to explain the main features of the observed spectra. Adding to these considerations the fact that the particle density within the current sheet and the magnetic field decrease as the wind expands to larger $R$, the conclusion is reached that the emitted luminosity, which depends on $N'_0$ and $B'_0$ will also decrease with radius, and this in combination with the decreasing peak of the spectrum means that the main contribution to the spectrum near the cutoff will still come from a very limited region close to $\\Rmin$. Therefore, possible changes in the dynamic evolution of the wind should not have a large effect to the observed spectrum and luminosity, while the exact shape of the current sheet should influence mainly the shape of the pulse, but not the cutoff energy or the emitted luminosity.  It is expected that reconnection will generally alter the physics of the current sheet beyond the pulsar's light cylinder, and might result in an evolution of the current sheet thickness and the wind Lorentz factor, especially in objects for which the radiative timescales of the thermal particles are very short in comparison to the evolution timescale of the magnetic field. Furthermore, particle acceleration takes place during the reconnection process which might result in a non-thermal tail to the thermal distribution, the characteristics of which will change as the current sheet evolves with radius. These phenomena will leave their imprint on the pulsed high-energy spectrum, and need to be investigated in a self-consistent way, something that is beyond the scope of this article. We can, however, give a simple order-of-magnitude argument about the reconnection-accelerated particles: deep in the current sheet particles can be accelerated by reconnection electric fields to energies higher than the thermal peak. These particles would have Lorentz factors extending to $\\gamma \\sim \\xi \\alc/\\Gamma$ (in the wind frame), thus giving rise to photons up to an energy \\begin{equation} \\mathcal{E}_{\\rm GeV}\\sim 8.8 \\frac{\\xi^2 \\Edot^{3/2}}{\\Gamma P} \\end{equation}where $\\xi = E'/B'_0 <1$. This could extend beyond the thermal peak for large enough $\\xi$ or $\\Edot$, whereas millisecond pulsars are again favoured by the dependence on the inverse of the period. If these high energy non-thermal particles escape the acceleration site and radiate in the field within the current sheet, their emission could give rise to a power-law tail extending beyond the GeV cutoff \\citep{zenitani+hoshino08}. This mechanism should be prominent for objects for which $t'_{\\rm s}/t'_{\\rm R}<1$, which tend to cluster at the upper left part of the $P-\\dot{P}$ diagram, as discussed. It is for these pulsars that acceleration by reconnection could become prominent and be observed in the form of power-law tails in the GeV-TeV regime.  Another possibility is that the particles in the current sheet are already accelerated to a power-law distribution when the sheet starts radiating, in which case a different distribution function would have to be used in order to calculate the current sheet parameters \\citep{balikhin08}. This might apply particularly to millisecond pulsars, which have lower surface magnetic fields which lead to lower pair production rates, and therefore less dense plasmas in their magnetospheres. In this case a non-thermal particle distribution seems likely to describe the physics of the current sheet in a more consistent way. We defer the investigation of such particle distributions to a future article. \\\\"
    },
    "1208/1208.2734_arXiv.txt": {
        "abstract": "We present the design, implementation and data taking performance of the MIcrowave Detection of Air Showers (MIDAS) experiment, a large field of view imaging telescope designed to detect microwave radiation from extensive air showers induced by ultra-high energy cosmic rays. This novel technique may bring a tenfold increase in detector duty cycle when compared to the standard fluorescence technique based on detection of ultraviolet photons. The MIDAS telescope consists of a 4.5 m diameter dish with a 53-pixel receiver camera, instrumented with feed horns operating in the  commercial extended C-Band (3.4 -- 4.2~GHz).  A self-trigger capability is implemented in the digital electronics. The main objectives of this first prototype of the MIDAS telescope - to validate the telescope design, and to demonstrate a large detector duty cycle  - were successfully accomplished in a dedicated data taking run at the University of Chicago campus prior to installation at the Pierre Auger Observatory. ",
        "introduction": "The origin and composition of Ultra-High Energy Cosmic Rays (UHECRs) remains uncertain~\\cite{LetessierSelvon:2011dy}, even after the progress made by the latest generation of experiments~\\cite{Abraham:2004dt,AbuZayyad:2000uu}. Due to the strong flux suppression above $10^{19}$~eV~\\cite{Abbasi:2007sv, Abraham:2008ru}, very large detection areas are necessary to study cosmic rays at these energies. A future UHECR Observatory based on standard techniques - Surface Detector arrays (SD) and Fluorescence Detectors (FD) -  may be  limited by cost and difficulty of deployment. In this context, radio detection techniques are attractive thanks to the low cost of individual elements, the little maintenance required and  a nearly 100\\% detection duty cycle.  Radio emission in the MHz range from extensive air showers (EAS) has been actively studied in the last decade by the LOPES~\\cite{Falcke:2005tc} and CODALEMA~\\cite{Ardouin:2005xm} experiments. MHz radio-detection is now well established, and AERA~\\cite{Fliescher:2012zz},  in commissioning phase at the Pierre Auger Observatory, will have a sufficiently large instrumented area to explore the potential of this technique for detection of the highest energy cosmic rays.  The use of the microwave (GHz) band for EAS detection was originally pursued by Jelley, Charman and collaborators~\\cite{Jelley:1966, Charman:1969} in the late 1960s, but abandoned due to the lack of satisfactory understanding of the emission mechanisms and to limitations in detector technology.  Thereafter it remained mostly unexplored, until recent laboratory measurements~\\cite{Gorham:2007af} with particle beams have renewed the interest in this part of the radio spectrum. These measurements suggest that microwave radiation is emitted from the weakly ionized air plasma of free electrons produced by the EAS induced ionization of the atmosphere. The radiation is expected to be continuous and relatively flat in frequency, unpolarized and emitted isotropically, and its intensity to scale with the number of particles of the shower.   \\begin{figure}[t] \\centering{ \\includegraphics[width=0.9\\columnwidth]{fig/telescope.jpg} } \\caption{The MIDAS telescope at the University of Chicago, with the 53-pixel camera at the prime focus of the 4.5~m diameter parabolic dish reflector.} \\label{fig:telescope} \\end{figure}  Detection of an isotropic emission in the GHz range  - akin to the detection of ultraviolet fluorescence photons by FD pioneered by the Fly's Eye experiment~\\cite{Baltrusaitis:1985mx}  and currently used by the Pierre Auger Observatory~\\cite{Abraham:2009pm} and the Telescope Array~\\cite{Tokuno:2012mi} - allows for the measurement of the EAS development in the atmosphere, which provides a calorimetric measurement of the energy and crucial information on the mass composition of primary cosmic rays. A GHz radio telescope would overcome the limitations of the FD technique, i.e. data taking only during moonless nights ($\\approx15\\%$ duty cycle) and significant systematic uncertainties introduced by light attenuation in the atmosphere. In fact, microwave detectors can operate 100\\% of the time, and attenuation in the GHz range is minimal, even with rain or clouds. Moreover, commercial off-the-shelf GHz equipment, mostly developed for satellite TV reception, is readily available and inexpensive. Microwave telescopes could provide existing UHECR experiments of unprecedented sensitivity to primary comic ray  composition, and be employed in a future large scale observatory.  Several complementary approaches to microwave detection of EAS are currently being pursued, including the AMBER and EASIER detectors~\\cite{Allison:2011zz} at the Pierre Auger Observatory, and the CROME experiment at KASCADE~\\cite{Smida:2011cv}. Also, new laboratory measurements with particle beams are being performed~\\cite{MAYBE,AMY} to better characterize the microwave emission.  In this paper, we present the  MIcrowave Detection of Air Showers (MIDAS) experiment, an imaging telescope whose primary objective is to  confirm the microwave emission from EAS, and to demonstrate the feasibility of a low cost design for this novel technique. The MIDAS telescope is described in  Sec.~2. The electromagnetic simulations of the telescope response are illustrated in Sec.~3. The calibration procedures and the measured sensitivity of the instrument are presented in Sec.~4.  A realistic simulation of the MIDAS detection of EAS is described in Sec.~5. Operation and data taking performance of the MIDAS telescope are presented in Sec.~6, and conclusions are drawn in Sec.~7.   \\begin{figure}[t] \\centering{ \\includegraphics[width=0.9\\columnwidth]{fig/camera_tr.jpg}} \\caption{A front view of the receiver camera, with LNBFs closely packed to maximize the sensitivity over the focal plane. } \\label{fig:camera} \\end{figure} ",
        "conclusions": "The MIDAS telescope - a prototype of an imaging detector for microwave emission from EAS - has been built and successfully operated at the University of Chicago. The telescope's design is based on inexpensive off-the-shelf microwave components, and a custom-made digital electronics and trigger system. The absolute calibration of the detector was established  with measurements of the Sun and other astrophysical objects as calibrated sources of microwave radiation. The sensitivity over the focal plane was determined with an EM simulation of the telescope validated by measurements of the Sun transit in the telescope's field of view. The sensitivity of the MIDAS detector to EAS has been studied with Monte Carlo simulations which include a realistic parameterization of the detector based on these calibration measurements. Several tens to several hundreds of events per year are expected for microwave flux intensities as suggested by laboratory measurements \\cite{Gorham:2007af}.  Several months of data taking in Chicago demonstrated that MIDAS can be reliably operated with minimal maintenance, and reach a duty cycle close to 100\\% even in an environment with high levels of RF background noise. Performances are expected to further improve in the radio quiet environment of the Pierre Auger Observatory, where MIDAS measurements in coincidence with the FD and SD will be essential to demonstrate the potential of this novel detection technique of UHECRs."
    },
    "1208/1208.5322_arXiv.txt": {
        "abstract": "{Young supernova remnants (SNRs) exhibit narrow filaments of non-thermal X-ray emission whose widths can be limited either by electron energy losses or damping of the magnetic field.}{We want to investigate whether or not different models of these filaments can be observationally tested.}{Using observational parameters of four historical remnants, we calculate the filament profiles and compare the spectra of the filaments with those of the total non-thermal emission. For that purpose, we solve an one-dimensional stationary transport equation for the isotropic differential number density of the electrons.}{We find that the difference between the spectra of filament and total non-thermal emission above 1~keV is more pronounced in the damping model than in the energy-loss model.}{A considerable damping of the magnetic field can result in an observable difference between the spectra of filament and total non-thermal emission, thus potentially permitting an observational discrimination between the energy-loss model and the damping model of the X-ray filaments.} ",
        "introduction": "Based on simple energetic considerations regarding the energy density of cosmic rays and the energy release per supernova explosion, SNRs have long been thought to be sources of galactic cosmic rays. This presumption is supported by numerous detections of non-thermal emission in the radio and X-ray band \\citep[e.g.,][]{1995Natur.378..255K,1999ApJ...525..357S,2001ApJ...548..814S,2000PASJ...52.1157B} observed from the direction of known SNRs and interpreted to be synchrotron radiation of relativistic electrons with energies up to $100\\,$TeV.  High-resolution observations of young SNRs performed with the \\emph{Chandra} satellite show that this emission of non-thermal radiation is concentrated in narrow regions on the limbs \\citep{2003ApJ...584..758V,2005ApJ...621..793B}. These regions of increased synchrotron emissivity close to the forward shock are called filaments and demonstrate the presence of high-energy electrons around their acceleration sites.  The most plausible process for the acceleration of electrons is diffusive shock acceleration (DSA), which leads to a power-law distribution of particles \\citep[e.g.,][]{1978MNRAS.182..147B,1978ApJ...221L..29B,1983RPPh...46..973D}. Although no clear evidence for relativistic-ion acceleration exists at shocks, the DSA-mechanism is also considered to be responsible for the acceleration of cosmic-ray nuclei, as indicated by observations of non-relativistic ion acceleration at solar-wind shocks driven by coronal mass ejections \\citep{2011ApJ...735....7R}. However, many details of the DSA are still vague such as the maximum energy of particles, the role of the magnetic field, and how the particles are injected into the acceleration process (also referred to as the injection problem). Apparently, the investigation of the properties of the non-thermal filaments may provide key information for a better understanding of the DSA-mechanism. In particular, knowing the magnetic-field strength gives constraints on the maximum particle energy achievable in the acceleration process, which can help answering the question whether SNRs can accelerate particles to energies above the knee in the cosmic-ray spectrum.  Accurate analysis of several SNRs shows that the filamentary structures are very thin compared with the radii of the remnants. This limitation of the filament widths is associated with a rapid decrease of the synchrotron emissivity that can be explained by energy losses of the electrons in a locally enhanced magnetic field. A number of authors have used that model to constrain the magnetic-field strength, the degree of turbulence and the obliquity \\citep[e.g.,][]{2003ApJ...589..827B,2006A&A...453..387P,2010ApJ...714..396A}. \\citet{2010ApJ...714..396A}, for instance, have investigated several filaments of the remnant of Cas A and found that the magnetic fields of the filaments are highly turbulent and nearly perpendicular to the shock normal. Another important result of this and other studies is an estimate of the downstream magnetic-field strength that is higher than simple shock compression would suggest. Such observations indicate an additional amplification of the magnetic field in the shock region of the SNRs. A possible amplification process could be a streaming instability in the upstream region as proposed by \\citet{2000MNRAS.314...65L} and \\citet{2004MNRAS.353..550B}, or the effects of preexisting turbulent density fluctuations on the propagating shock front \\citep{2007ApJ...663L..41G}.  Besides energy losses, also the magnetic field itself can limit the filament widths. Based on the turbulence relaxation downstream of the forward shock and neglecting any amplification process, \\citet{2005ApJ...626L.101P} have calculated that the turbulent magnetic field downstream can decay exponentially on a damping-length scale $l_{\\mbox{\\tiny d}}=10^{16}-10^{17}\\,$cm that is small enough to produce the narrow observable filaments. Furthermore, from observations of the post-shock steepening of the synchrotron spectrum in Tycho's SNR it can be seen that also the damping of the magnetic field fairly well describes the corresponding X-ray data \\citep{2007ApJ...665..315C}, and thus, may appear within the filaments.  Since the magnetic field controls the radiative cooling of electrons, high magnetic fields lead to strong cooling, and too few high-energy electrons remain capable of producing the gamma-ray emission, that is observed from regions near the edges of numerous SNRs \\citep[e.g.,][]{2007ApJ...661..236A,2010A&A...516A..62A,2011ApJ...734...28A}. Any gamma rays observed in such a case are likely to be hadronic in origin. Weak cooling leads to a large number of high-energy electrons and the possibility of gamma-ray emission through inverse Compton or bremsstrahlung processes. All these implications of the magnetic-field structure on the particle acceleration, gamma-ray emission and magnetic-field amplification make it necessary to understand the non-thermal filaments in detail.  In this paper we investigate the properties of the filaments for both cases, filaments limited by electron energy losses or by damping of the magnetic field. For that purpose, using observational values of some characteristical SNR parameter, we calculate the X-ray emission of the filaments. The resulting filament profiles then allow us to make specific predictions regarding the magnetic-field strength. We additionally calculate the total non-thermal emission, which shall be referred to as \"plateau\", and compare their spectra with those of the filaments. It should be noted that in our models we only consider non-thermal synchrotron emission and restrict ourselves to the evolution of relativistic electrons in the downstream region. Furthermore, we assume the electrons to be already accelerated at the shock front and treat our problem to be independent of the acceleration process. We also do not consider any electron propagation into the upstream region and simplify the SNRs to be spherical objects of constant downstream-velocity profile. Recent hydrodynamical simulations suggest that this oversimplification of a constant velocity is an acceptable approximation only for SNRs of an age of less than several hundred years \\citep{2012APh....35..300T}, implying that our models are restricted to SNRs being in the adiabatic phase and just entering the Sedov phase, respectively. ",
        "conclusions": "Compared to the magnetic-field damping model, the spectra of filament and plateau obtained in the energy-loss model exhibit larger spectral indices. This can be explained by the considerable energy losses leading to the evolution towards a softer electron distribution in the energy-loss model, and hence, resulting in X-ray spectra that are softer than those obtained in the damping model.  In case of a weak magnetic-field damping the difference between the spectral indices of filament and plateau over the full X-ray spectrum is smaller than 0.1, which is probably to small to be detectable. Only if there is a strong damping, our calculation suggests a measurable difference between the spectra in some SNRs, since the difference between the spectral indices of filament and plateau can take values of almost 0.2 at X-ray energies higher than 1~keV.  In the energy-loss model the difference between the indices of filament and plateau above the X-ray energy of 1~keV is even smaller than 0.1, so that a possible detection can be excluded here, too. On the other hand, the difference between the indices below 1~keV is larger than 0.1, and at a photon energy of 0.1~keV it is even approximately 0.3. This might suggest that there is a measurable difference in the spectra of filament and plateau at small X-ray energies, if the filaments are limited by energy losses. However, on account of the interstellar photoelectric absorption of the soft X-rays, these different spectral characteristics are probably not detectable, too. Furthermore, the plasma downstream of the forward shock is at high temperature, implying also thermal emission contributing to the soft X-ray band, and thus, complicating a clear identification of the non-thermal emission.  Hence, if there is no measurable difference between the spectra of filament and plateau, it is not possible to make definite predictions from the comparison of the spectra whether the filaments are limited by energy losses of the radiating electrons or by damping of the magnetic field. But if a significant difference appears, our calculations then suggest that the filaments are limited by the magnetic field itself.  It should be noted that our results presented here have been derived using Bohm diffusion. According to Eq. (\\ref{eqn:diffusion length}), a larger diffusion coefficient with gyrofactor $\\eta>1$ would imply a larger diffusion length, resulting in significant widening of the filaments, because now, the regions farther from the shock contain a sufficient number of high-energy electrons contributing to the intensity. Widening must then be compensated by a higher magnetic field to retain the observed filament widths. Hence, the magnetic-field strengths derived in our models represent lower limits for the chosen parameters. The calculation then shows that a larger diffusion coefficient results in softer spectra due to a lower cut-off energy, which decreases with increasing gyrofactor. However, the final results regarding the differences in spectral indices do not change fundamentally.  In a last step we want to compare the predictions derived here with observations. At first, we notice that, independently of the model, the parameters from the remnants of Cas A, Tycho, Kepler and SN 1006 lead to nearly the same spectral behaviour in case of similar shock velocities, as can be seen from the spectral indices in Table \\ref{tab:Constraints energy-loss model} and Table \\ref{tab:Constraints magnetic damping model}. However, the analysis of the filament spectra of these remnants reported by \\citet{2003ApJ...589..827B,2005ApJ...621..793B} reveals significant differences among the spectral indices obtained from the fit of an absorbed power-law model. Compared to our spectra whose calculation has been done using an injection index resulting from an unmodified shock ($s=2$), the observation may be an indication for different electron injection indices in these remnants, implying shocks that are differently affected by non-linear effects due to differences in efficiency in the particle acceleration.  Regarding the magnetic-field strengths, we take, as an example for comparison with our results, the non-thermal filaments of Cas A analysed by \\citet{2010ApJ...714..396A}. From the best-fit parameters used to fit the observed filament spectra, the magnetic field has been derived to be in the range $(30-70)~\\mu$G. These values are consistent with those derived from the magnetic-field damping model, in which the magnetic field varies, according to Eq. (\\ref{eqn:magnetic field}) and the values from Table \\ref{tab:Constraints magnetic damping model}, between the field strengths $(10-260)~\\mu$G for weak damping and $(10-115)~\\mu$G for strong damping, respectively. For comparison, the constant magnetic field derived from the energy-loss model is several times higher, $B=520~\\mu$G. This might suggest that the non-thermal filaments of Cas A are limited by the damping of the magnetic field.  Another comparison concerns the magnetic fields in SN 1006 and Tycho. Using the data from radio up to TeV-observations, \\citet{2010A&A...516A..62A} have analysed the multi-wavelength spectrum of SN 1006 in the framework of a leptonic and hadronic origin for the gamma-ray emission, giving a magnetic-field of $\\sim 30~\\mu$G in the leptonic model and a magnetic field of $\\sim 120~\\mu$G in the hadronic model, respectively. Moreover, combining radio and X-ray data with recent TeV-observations performed with the VERITAS instrument, the magnetic field of Tycho has been estimated to be $\\sim 80~\\mu$G in a leptonic-dominated model, whereas a hadronic dominated model yields a magnetic field of $\\sim 230~\\mu$G \\citep{2011ApJ...730L..20A}. Compared to our model predictions given in Table \\ref{tab:Constraints energy-loss model} and Table \\ref{tab:Constraints magnetic damping model}, we notice that the magnetic fields derived from the energy-loss model are in good agreement with those estimated from the hadronic model used to describe the observed spectra of SN 1006 and Tycho. In contrast, the predictions from the magnetic-field damping model suggest the leptonic model for the origin of the gamma-ray emission from these remnants. It should be noted that current gamma-ray observations do not reach the spatial resolution of those done in X-rays, so that the magnetic fields estimated using gamma-ray observations of SN 1006 and Tycho are averages over a region much larger than the filaments, implying that the observed values do not necessarily match those found for the filaments.  To discriminate between the energy-loss model and magnetic-field damping model, and hence between a leptonic and a hadronic origin of TeV-band gamma-ray emission, one may either search for differences between X-ray spectra of filaments and plateau, as calculated in this paper, or perform gamma-ray observations with higher spatial resolution.  \\subsection*{Acknowledgement}  We acknowledge support by the \"Helmholtz Alliance for Astroparticle Phyics HAP\" funded by the Initiative and Networking Fund of the Helmholtz Association."
    },
    "1208/1208.1117_arXiv.txt": {
        "abstract": "{  In this work we analyze the spectroscopic properties of a large number of \\ion{H}{ii} regions, $\\sim$2600, located in 38 galaxies. The sample of galaxies has been assembled from the face-on spirals in the PINGS survey and a sample described in M\\'armol-Queralt\\'o (2011, henceforth Paper I). All the galaxies were observed using Integral Field Spectroscopy with a similar setup, covering their optical extension up to $\\sim$2.4 effective radii within a wavelength range from $\\sim$3700 to $\\sim$6900\\AA.  We develop a new automatic procedure to detect \\ion{H}{ii} regions, based on the contrast of the H$\\alpha$ intensity maps extracted from the datacubes. Once detected, the procedure provides us with the integrated spectra of each individual segmented region. In total, we derive good quality spectroscopic information for $\\sim$2600 independent \\ion{H}{ii} regions/complexes. This is by far the largest nearby 2-dimensional spectroscopic survey presented on this kind of regions up-to-date. Even more, our selection criteria and dataset guarantee that we cover the regions in an unbiased way, regarding the spatial sampling.  A well-tested automatic decoupling procedure has been applied to remove the underlying stellar population, deriving the main properties (intensity, dispersion and velocity) of the strongest emission lines in the considered wavelength range (covering from [\\ion{O}{ii}]~$\\lambda$3727 to [\\ion{S}{ii}]~$\\lambda$6731). A final catalogue of the spectroscopic properties of these regions has been created for each galaxy. Additional information regarding the morphology, spiral structure, gas kinematics, and surface brightness of the underlying stellar population has been included in each catalogue.  In the current study we focused on the understanding of the average properties of the \\ion{H}{ii} regions and their radial distributions. We found that, statistically, there is a significant change of the ionization conditions across the optical extent of the galaxies. The fraction of \\ion{H}{ii} regions located in the intermediate ionization range in a classical BPT diagram is larger for the central regions ($r< 0.25 r_e$), than in the outer ones. This is somehow expected, if the origin of this shift is due to the contamination of non-starforming ionization sources (e.g., AGN, Shocks, post-AGB stars, etc.), that occur more frequently in the center of the galaxies.   We find that the gas-phase oxygen abundance and the H$\\alpha$ equivalent width present negative and positive gradient, respectively. The distribution of slopes is statistically compatible with a random Gaussian distribution around the mean value, if the radial distances are measured in units of the respective effective radius. No difference in the slope is found for galaxies of different morphologies: barred/non-barred, grand-design/flocculent. Therefore, the effective radius is a universal scale length for gradients in the evolution of galaxies.  Other properties have a larger variance across each object, and galaxy by galaxy (like the electron density), without a clear characteristic value, or they are well described by the average value either galaxy by galaxy or among the different galaxies (like the dust attenuation).  } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.3781_arXiv.txt": {
        "abstract": "Spectra of Supernovae of type Ia (\\snia) are commonly interpreted as a continuum with absorption features. The pseudo equivalent width (pEW) and Doppler shift of absorption features like \\SitAbs, \\Cat\\ HK 3750\\AA, and the \\Cat\\ IR-triplet 8150\\AA\\, measures which are commonly interpreted as the optical thickness and the velocity of a shell of corresponding absorbing material, form the basis of common \\snia\\ spectral interpretation and classification schemes.  In this paper, we examine the nature of spectral features in \\snia\\ spectra using W7 model spectra, and show that \\SitAbs\\ and many other features are largely emission dominated instead of absorption dominated. We show that apparent absorption features (like \\SitAbs) are frequently just coincidental troughs between two (or more) uncorrelated emission features. The pEW measured between such emission peaks is then only little related to true strength of the presumed absorption feature. It shows how the concepts of ``absorption troughs'' and ``continuum'' can be misleading and should be used with care. Furthermore, using the same model spectra, we demonstrate for different times post explosion \\emph{how} spectral features overlap each other and how together they compose the total spectrum. This overlap distorts individual line profiles and affects measured absorption velocities.  With an explosion model atmosphere that is tuned to match specific observations, the method presented in this paper can in principle be used to quantify all these effects and improve the interpretation and informative value of observed \\snia\\ spectra. ",
        "introduction": "Supernovae of Type Ia (\\snia) generally show very similar spectral and photometric behavior (e.g., \\citealt{Filippenko97}). Because of this interesting homogeneity (and the fact that these objects become extremely bright) \\sneia\\ have proven to be very useful objects to study the universe at large distances \\citep{Riess98, Perlmutter99}. Yet despite the general similarity, there is an even more interesting diversity among \\sneia, both photometrical and spectral. Whereas photometrical diversity has long been studied \\citep{Phillips93, Wang03}, good spectroscopic data have only recently become available in large numbers (CfA SN Program \\citealt{Blondin12}; Berkeley \\snia\\ Program \\citealt{Silverman12}).  One step towards understanding spectroscopic diversity is finding similarities and defining classes. A number of spectral classification schemes for \\snia\\ have been proposed in literature. These are based on measurements performed on conspicuous \\Sit\\ and \\Cat\\ absorption features present in all ``normal'' \\citep{Branch93} \\snia. These absorption feature measurements are: 1) The feature strength, expressed as pseudo equivalent width (pEW, see \\citealt{Folatelli04}) \\citep{Branch06}; 2) The line blue-shift, expressed as expansion velocity \\citep{Wang09}; 3) The rate of change in blue-shift, expressed as velocity gradient \\citep{Benetti05}.  It is well known that atomic transition lines formed in a rapidly expanding atmosphere, like \\snia, form P-Cyg profiles \\citep{Lamers99}. P-Cyg profiles generally show a superposition of an emission wing that is more or less symmetric around the transition's rest wavelength plus an absorption wing that is blue shifted. But the precise shape of P-Cyg profiles in relativistically expanding atmospheres, and specifically the ratio of emission to absorption, strongly depends on details of the atmosphere \\citep{Hutsemekers93}: 1) The radial extension of the line forming region; 2) The velocity gradient across the line forming region; 3) The wavelength dependence of the ``continuum'' and its change across the line forming region.  Parameterized models of \\snia\\ atmospheres, like SYNOW \\citep{Fisher00, Branch09} or SYNAPPS \\citep{Thomas11}, have proven to be very powerful and rewarding tools for the interpretation of \\snia\\ spectra. Specifically, these models made it possible to identify lines and study the composition of observed \\snia. However, these models do not account for the geometrical extension of the line forming regions nor the velocity gradients and change of the ``continuum'' flux across the line forming regions involved in shaping the P-Cyg profiles. Therefore, such models can not be expected to give accurate and physically realistic profile shapes, even though the large number of free parameters usually enables a rather precise reproduction of observed spectra.  Another fact that complicates the interpretation of \\snia\\ spectra is the large number of transition lines that leave their signature in the spectra in an indistinguishable way (i.e., smeared out by the large expansion velocities of \\snia\\ atmospheres). \\cite{Kasen06b} have shown that many millions of transitions are needed to calculate the \\snia\\ opacity accurately. Even though some of these lines are stronger than others, the background opacity formed by the millions of weaker lines is definitely not constant (see \\citealt{Kasen06b}) and will leave its signature on the profiles of those stronger lines.  Clearly, the line profiles in \\snia\\ spectra contain a large amount of information about the detailed atmospheric conditions and understanding their nature is the key to elevate the use and yields of spectroscopic observations to a higher level. In this paper we will demonstrate that \\SitAbs\\ and many other features are in fact emission dominated and that commonly used measures like the pEW and the absorption velocity are little related to the actual properties of the these features. ",
        "conclusions": "\\subsection{Absence of an important emission peak} Generally, W7 provides a reasonably good fit to the Hsiao07 templates. There is, however, one \\emph{very serious mismatch}. The model does not describe the 5800\\AA\\ emission peak observed on day 20, 30, and 40pe. (Note that this finding is in agreement with results from W7 calculated with other radiation transport codes (e.g., \\citealt{Kasen06, Kromer09}).) This is a serious problem since this emission peak is responsible for the blue side of the \\SitAbs\\ trough and thus affects the \\Sit\\ pEW measurements.  What makes this problem even more serious is that other \\snia\\ explosion models do not describe this emission peak either (e.g., \\citealt[Figure 2b]{Kasen09}; \\citealt[Figure 3]{Roepke12}). There are two possible reasons for this emission peak to be absent from the models: 1) the chemical element(s) responsible for this observed peak is under-represented in the explosion models, 2) the opacity data for the elements present in the models is missing one or more important transition lines\\footnote{ In the scope of this work, we have verified the opacity data used by \\phx\\ against recent atomic data from R.L. Kurucz (http://kurucz.harvard.edu/). }. Note that NLTE effects do affect the shape of \\snia\\ spectra (\\citealt{Baron96}; van Rossum, in preparation) but are unlikely to explain discrepancies as big as this missing emission peak.  Since the origin of this observed emission line is unknown and it does shape the blue side of the \\SitAbs\\ absorption trough, special care must be taken with the interpretation of the pEW of this famous trough.   \\subsection{Blended features} Line features are known to strongly blend in \\snia\\ spectra but still it is common to speak of \\emph{identified} features, named after the predominant contributor. Figure \\ref{fig:0lin_all} and \\ref{fig:0lin_diffeach} show the degree of blending in the model spectra, and the way it changes over time, for each feature in its own way. Note that the detailed degree of blending depends on the atmospheric conditions, which vary between \\snia\\ objects. The W7 results presented here provide a qualitative picture of the blending effects.% A \\snia\\ model atmosphere tuned to fit a certain observation would allow to quantify the degrees of blending using the Delta method presented in this paper.% This would be a very effective way to \\emph{increase the informative value of observed \\snia\\ spectra}.%   \\subsection{The \\SitAbs\\ feature} Since this is the most widely used feature in the characterization and classification of \\snia\\ spectra it deserves extra discussion.  The absorption velocity of the \\SitLab\\ line versus the 6150\\AA\\ trough is demonstrated in Figure \\ref{fig:si2_v0}. It compares the velocity of isolated Si with the value measured in the full spectrum. At early times, the \\SitAbs\\ trough is filled from the blue side by emission from the next bluer \\Sit\\ line leading to lower velocities, although the effect is small (a few percent).  Figure \\ref{fig:si2_ew} shows that from day 26pe on the \\Sit\\ emission gradually declines. At the same time the peak is seamlessly taken over by Fe\\,{\\sc ii} emission at a similar but slightly higher wavelength. This not only affects the pEW but also makes the \\Sit\\ velocity at late times appear lower (and its gradient higher) than it really is.  The run of the full spectrum curve in Figure \\ref{fig:si2_ew} does not match the observed, gradually declining behavior (e.g., \\citealt[Figure 15]{Blondin12}). Also, the 5800\\AA\\ emission feature, that is missing in the model (see previous subsection), causes the \\Sit\\ absorption trough to disappear after day 28pe. This missing feature apparently plays \\emph{a key role} in run of the absorption velocity.  For these reasons, using the \\SitAbs\\ velocity at late times (e.g., \\citealt{Howell06, Blondin12}) is potentially problematic. The same problems might affect the velocity gradient definition in \\cite{Benetti05}. The ``improved'' velocity gradient definition of \\cite{Blondin12} limits using \\SitAbs\\ velocities to day 10 post maximum brightness (approx. day 28pe for W7) and thus partly avoids these problems.    \\subsection{The nature of spectral features in \\snia\\ spectra} In this subsection we explain how the concepts of ``absorption troughs'' and ``continuum'' can be misleading and should be used with care.%  \\paragraph{Misinterpretation of ``absorption troughs''} The \\emph{identified} features commonly used for interpretation and classification of \\snia\\ spectra describe the apparent ``absorption troughs'', with apparent ``continuum'' in between of those troughs. However, many of the features in \\snia\\ spectra are predominantly in emission. Figure \\ref{fig:0lin_diffeach} shows that for wavelengths larger than 5400\\AA\\ the spectral features generally are emission dominated (except possibly \\Sit\\ at very early times), and on day 20pe (close to peak luminosity, day 18pe) emission contributes significantly to the flux level down to wavelengths as low as 3800\\AA. The widths and depths of these ``absorption troughs'' largely depend on the strengths and separation of the enclosing emission features.  \\paragraph{Misinterpretation of ``continuum flux level''} Estimating a realistic pseudo continuum from a \\snia\\ spectrum is not trivial. But since line emission is strong, a ``maximum'' continuum level that connects observed flux peaks is definitely unrealistic for many features at most evolution times.  \\paragraph{Redefinition of \\snia\\ feature strength} A \\snia\\ model atmosphere tuned to fit a certain observation would in principle allow to determine the pseudo continuum for any spectral feature using the method presented in this paper.% Yet such atmospheres are not generally available.%  Therefore, instead of the common pEW method, we suggest to measure the strength of spectral features using the height difference of P-Cyg profiles (at least in the region where emission is strong in \\snia\\ spectra, see above). The stenght is defined as difference in flux levels at the trough minimum and the peak maximum on the red side of the trough normalized to the mean of these levels. The emission peak on the blue side of the trough, the strength of which is unrelated to the strength of the feature of interest, can thus be excluded, potentially giving a more accurate strength measure. Note that this new definition is a dimensionless quantity, not taking into account the width of the feature, which is loosely related to the absorption velocity. Thus the strength and the velocity measures become potentially more independent. Application of this redefined strength measure to large \\snia\\ spectral datasets will show whether this measure improves correlations to other spectroscopic or photometric properties, and how it will affect classification schemes based on line strengths.  In future work we will take the method presented in this paper to a higher level by using model atmospheres that are tuned to match observed \\snia\\ spectra. This will be parametrized atmospheres, based on W7 and more modern explosion models. This way, the details of the atmospheric structure can be directly inferred from observed \\snia\\ spectra in a physically self-consistent way.  \\hspace{.5cm}"
    },
    "1208/1208.4959_arXiv.txt": {
        "abstract": "{Planets are thought to eventually form from the mostly gaseous ($\\sim 99\\%$ of the mass) disks around young stars. The density structure and chemical composition of protoplanetary disks are affected by the incident radiation field at optical, far-ultraviolet (FUV), and X-ray wavelengths, as well as by the dust properties.}{The effect of FUV and X-rays on the disk structure and the gas chemical composition are investigated. This work forms the basis of a second paper, which discusses the impact on diagnostic lines of, e.g., C$^+$, O, H$_2$O, and Ne$^+$ observed with facilities such as Spitzer and Herschel.}{A grid of 240 models is computed in which the X-ray and FUV luminosity, minimum grain size, dust size distribution, and surface density distribution are varied in a systematic way. The hydrostatic structure and the thermo-chemical structure are calculated using \\underline{Pro}toplanetary \\underline{Di}sk \\underline{Mo}del (ProDiMo), with the recent addition of X-rays.}{The abundance structure of neutral oxygen is very stable to changes in the X-ray and FUV luminosity, and the emission lines will thus be useful tracers of the disk mass and temperature. The C$^+$ abundance distribution is sensitive to both X-rays and FUV. The radial column density profile shows two peaks, one at the inner rim and a second one at a radius $r=5-10$~AU. Ne$^+$ and other heavy elements have a very strong response to X-rays, and the column density in the inner disk increases by two orders of magnitude from the lowest ($L_{\\rm X} = 10^{29}$~erg\\,s$^{-1}$) to the highest considered X-ray flux ($L_{\\rm X} = 10^{32}$~erg\\,s$^{-1}$). FUV confines the Ne$^+$ ionized region to areas closer to the star at low X-ray luminosities ($L_X = 10^{29}$~erg\\,s$^{-1}$). H$_2$O abundances are enhanced by X-rays due to higher temperatures in the inner disk than in the FUV only case, thus leading to a more efficient neutral-neutral formation channel. Also, the higher ionization fraction provides an ion-molecule route in the outer disk. The line fluxes and profiles are affected by the effects on these species, thus providing diagnostic value in the study of FUV and X-ray irradiated disks around T Tauri stars.}{} ",
        "introduction": "New observing facilities in the past decade pushed our understanding of protoplanetary disks from a rough picture of a vertically layered structure to a wealth of details on the composition and two-dimensional structure of the gas and dust of these disks. In the infrared, the Spitzer Space Telescope performed systematic studies of nearby star-forming regions. The spectral energy distributions (SEDs) revealed the physical structure of disks, e.g., the presence of gaps, source-to-source variations and important statistics on SED types, which allow evolutionary scenarios to be built \\citep[cf.][]{Merin2010}. Extensive gas-phase emission line studies with Spitzer provide a first indication of chemical diversity across stellar spectral types \\citep[cf.][]{Najita2009,Pontoppidan2010,Lahuis2007}. In the near-infared, there have been large ground-based studies with, e.g.,\\ the VLT and the Keck telescope resolving the lines to study the kinematics (spatial origin of the lines in hot disk surfaces and winds) and their excitation mechanisms \\citep[thermal and fluorescence; ][]{Brittain2010, Pontoppidan2011, Fedele2011}. In the past two years, the Herschel Space Observatory extended the spectral window to the far-IR and submm, with spectral line scans within the Dust, Ice, and Gas in Time (DIGIT) key program (PI: N. Evans) for many Herbig disks, e.g., HD~100549 \\citep{Sturm2010}, water studies of a few selected targets, such as TW Hya \\citep{Hogerheijde2011} from the Water in Star-forming Regions with Herschel (WISH) key program (PI: E.van Dishoeck), and large statistical gas surveys targeting the dominant cooling lines \\citep[e.g.,][]{Mathews2010,Meeus2012,Dent2012, Riviere-Marichalar2012} in the Gas in Protoplanetary Systems (GASPS) key program (PI: B. Dent). The HIFI and some PACS line detections deal with the cold water, while the 63.3~$\\mu$m H$_2$O line possibly probes the inner water reservoir.  Observational studies report trends in emission line strength with the irradiation of the central star. \\citet{Guedel2010} find that the [Ne\\,{\\sc ii}] emission at 12.81\\,$\\mu$m correlates with the X-ray luminosity of the central star; the slope of the correlation for nonjet sources is $0.44\\pm 0.18$. \\citet{Pinte2010} and \\citet{Meeus2012} show that the [O\\,{\\sc i}]\\,63\\,$\\mu$m line flux increases with stellar luminosity and that this is largely driven by the strength of the far-ultraviolet (FUV) luminosity. \\citet{Riviere-Marichalar2012} report a tentative trend of the 63.3\\,$\\mu$m water line flux with X-ray luminosity.  It is obvious from previous works that the most relevant parameter for the SED is the total stellar luminosity; in most cases, $L_{\\rm X}$ and $L_{\\rm FUV}$ present only a small fraction of this and hence by themselves do not cause significant SED changes. However, recent near-IR high-contrast imaging with HiCIAO on the Subaru 8.2\\,m telescope allowed probing of the inner disk structure on scales below 0.1\\arcsec\\ \\citep{Thalmann2010,Hashimoto2011}. This offers the possibility of a direct probe, e.g., the height of the inner rim as predicted by protoplanetary disk models. \\citet{Thi2011} predicted that the height of the inner rim could be a factor two higher when the vertical hydrostatic structure takes into account the gas temperature at the inner rim. \\citet{Aresu2011} show that the height of the inner rim increases in disk models with the impinging stellar $L_{\\rm X}$.  In the last decade, much theoretical progress has taken place to support the interpretation of the wealth of new observations, specifically the gas observations. The foundation was laid by \\citet{Chiang1997}, \\citet{D'Alessio1998}, and \\citet{Dullemond2001}, whose layered dust disk models have helped to interpret SEDs for irradiated disks (see Dullemond et al. \\citeyear{Dullemond2007} for a review of disk structure models). Based on this, several groups focussed on the effects of FUV and X-rays on the thermal and chemical structure of the gas in a pre-prescribed protoplanetary disk model, e.g., \\citet{vanZadelhoff2003}, \\citet{Kamp2004}, \\citet{Jonkheid2004}, and \\citet{Glassgold2004}. In a subsequent step, \\citet{Nomura2005}, \\citet{Gorti2008}, and \\citet{Woitke2009} studied the chemical structure of disks, while self-consistently solving for the vertical hydrostatic equilibrium using the gas temperature. Most recently, \\citet{Aresu2011} performed an exploratory study to assess the relative importance of FUV and X-rays by expanding the existing \\underline{Pro}toplanetary \\underline{Di}sk \\underline{Mo}del (ProDiMo) code to include X-ray processes. The modeling efforts that solve for the vertical disk structure are computationally expensive because the chemical networks, heating/cooling balance, 2D continuum radiative transfer, and hydrostatic equilibrium have to be solved iteratively. Hence, these studies have largely focussed on a single representative disk model, or at most a handful of models varying one specific parameter. However, large-model grids are in principle required to understand the potential diagnostic power of certain observables. So far examples for this are large grids (of the order of 200\\,000-300\\,000 models) of parametrized dust disk models to study SED diagnostics \\citep{Robitaille2006} and of parametrized gas/dust disks to study the gas emission line diagnostics \\citep{Woitke2011,Kamp2011}.  In this paper, we perform for the first time an extensive analysis of a grid of 240 self-consistent disk models (including the vertical disk structure) to study the effects of X-rays, FUV, and the relative importance of grain size and gas surface density distribution on the thermal, chemical, and physical structure of disks around T Tauri stars. This paper builds on the implementation of X-rays into the ProDiMo code as described in \\citet{Aresu2011}. While we discuss here foremost the physical and chemical structure of the models and how they change with irradiation, a companion paper \\citep[][paper II]{Aresu2012}, discusses the power of line diagnostic in disentangling these effects. The paper is structured in the following way: Sect.~\\ref{updates} describes the updates on the code and the range of parameters used in the grid. The effects on disk temperature and density structure will be discussed in Sect.~\\ref{dens_temp}. The resulting distribution of various key species and its key reactions are extensively described in Sect.~\\ref{chemistry}. Sect.~\\ref{rad_col} shows the radial column density distributions. Those are key in understanding the line profiles, which is the topic of the accompanying paper II. Sect. \\ref{conclusions} summarizes the conclusions and implications of the paper. ",
        "conclusions": "In this paper, we discussed the combined effects of FUV and X-rays on both the density structure and the thermal and chemical balance of disks around T Tauri stars for an expected range of parameters (dust size distribution, density profile, etc.), yielding a total of 240 models. Here we highlight the main results and a few implications:  \\subsection{The disk thermal and chemical structure}  {\\it Temperature structure:} The extent of the disk where the temperature is higher than $T > 1000$~K is much larger when X-rays are included. X-rays have a much higher heating efficiency than FUV, $30 - 50$\\% compared to $< 3$\\%, respectively.  {\\it Density structure:} Increasing FUV luminosities does not change the scale height of the inner rim; it only alters the width and height of the second bump in the disk that is created at intermediate radii ($r \\sim 3 - 10$~AU), behind the region shielded by the puffed-up inner rim. Gas temperatures at the inner rim are much higher when X-rays are included and, as a result, the inner rim is puffed up to higher and higher altitudes for increasing X-ray luminosities. As a direct extension of this theoretical work, the existence of the second bump could potentially be tested by continuum imaging face-on protoplanetary disks in the near-infrared with, e.g., VLT or Keck.  {\\it Scale height:} Considering only FUV, we see that the scale height shows a maximum in the unattenuated parts ($z/r > 0.5$) of the disk. When X-rays are added we find that this maximum is smoothed over a larger region (out to $r \\sim 10$~AU). The scale height in these regions is larger than one would expect from the flaring index in the outer regions of the disk. When moving to smaller relative height, e.g, $z/r = 0.1$, the break in the flaring index disappears. Another observational possibility would be to do continuum interferometry in the near-IR (VLTI) to directly measure the physical height of the inner rim.  \\subsection{Chemical balance}  {\\it Ionization fraction:} The ionization fraction reaches values as high as $x_{\\rm e^-} \\sim 10^{-2}$ in exceptional cases in our FUV-only models, whereas X-rays easily maintain these ionization fractions throughout large portions of the disk. Even when the gas becomes partially shielded, it can still maintain a significant ionization fraction, leading to an ion-molecule chemistry that can form molecules at low temperatures, which is not possible with neutral-neutral reactions as they usually have temperature barriers.  {\\it Formation of H$_2$ through the H$^-$ route:} Formation of H$_2$ dust is usually much more efficient in environments that have solar metallicities. Because of the high ionization fraction due to X-ray irradiation, the formation route $\\rm H^- + H \\rightarrow H_2 + e^-$ is able to provide a significant addition of order 50\\% percent to the H$_2$ on dust formation route. Overall, the H$_2$ to H abundance ratio is increased by at least two orders of magnitude when X-rays are present.  {\\it Formation of water and OH:} The OH and H$_2$O abundances are more concentrated toward the inner regions of the disk, when only FUV is irradiating the disk. This is because the temperature is only there sufficient to drive the neutral-neutral formation route. The outer disk shows significant enhancements (up to two orders of magnitude) when X-rays are added. The higher ionization fractions make it possible to form the species through ion-molecule reactions and sustain abundance levels of $x_{\\rm H_2O}\\sim 10^{-6} - 10^{-7}$. Such abundance levels cannot be reached in outer disk models without X-rays. Only the warm inner disks allow even higher levels of water abundance through warm neutral-neutral chemistry.   {\\it Resulting abundance structures and radial column density profiles:} Whereas neutral oxygen and CO are very stable to both FUV and X-rays, this is not the case for, e.g., C$^+$, Ne$^+$, and H$_2$O. Ne$^+$ is strongly enhanced by X-rays and confined to the inner regions by larger FUV luminosities. This latter aspect will certainly affect the line widths. A high $L_{\\rm X}/L_{\\rm FUV}$ ratio is favorable for water formation, especially in the outer disk.  \\subsection{Outlook}  In paper II, we will perform a radiation transfer analysis of the aforementioned species and correlate line fluxes and line widths directly to the FUV and X-ray luminosities. This will allow a discussion of the diagnostic value of these species and provide a theoretical framework for the interpretation of observational data. We will also discuss our results in the context of data obtained within several observational efforts (Spitzer, Herschel, and ground-based observing programs)."
    },
    "1208/1208.5028_arXiv.txt": {
        "abstract": "One significant difference between the atmospheres of stars and exoplanets is the presence of condensed particles (clouds or hazes) in the atmosphere of the latter. In current 1D models clouds and hazes are treated in an approximate way by raising the surface albedo, or adopting measured Earth cloud properties. The former method introduces errors to the modeled spectra of the exoplanet, as clouds shield the lower atmosphere and thus modify the spectral features. The latter method works only for an exact Earth-analog, but it is challenging to extend to other planets.  The main goal of this paper is to develop a self-consistent microphysical cloud model for 1D atmospheric codes, which can reproduce some observed properties of Earth, such as the average albedo, surface temperature, and global energy budget. The cloud model is designed to be computationally efficient, simple to implement, and applicable for a wide range of atmospheric parameters for planets in the habitable zone.  We use a 1D, cloud-free, radiative-convective, and photochemical equilibrium code originally developed by Kasting, Pavlov, Segura, and collaborators as basis for our cloudy atmosphere model. The cloud model is based on models used by the meteorology community for Earth's clouds. The free parameters of the model are the relative humidity and number density of condensation nuclei, and the precipitation efficiency. In a 1D model, the cloud coverage cannot be self-consistently determined, thus we treat it as a free parameter.  We apply this model to Earth (aerosol number density 100 cm$^{-3}$, relative humidity 77 \\%, liquid cloud fraction 40\\%, and ice cloud fraction 25\\%) and find that a precipitation efficiency of 0.8 is needed to reproduce the albedo, average surface temperature and global energy budget of Earth. We perform simulations to determine how the albedo and the climate of a planet is influenced by the free parameters of the cloud model. We find that the planetary climate is most sensitive to changes in the liquid water cloud fraction and precipitation efficiency.  The advantage of our cloud model is that the cloud height and the droplet sizes are self-consistently calculated, both of which influence the climate and albedo of exoplanets. ",
        "introduction": "Clouds on exoplanets influence the temperature-pressure profile (T-P profile) of the atmosphere owing to their high opacity and therefore influence the borders of the Habitable Zone \\citep[see e.g.,][]{Forget1997, Selsis2007, Kitzmann2010}. They also influence the detectable spectra by reducing the observable atmospheric features \\citep[see e.g.,][]{Kaltenegger2007,Kitzmann2011}. Clouds are also a source of uncertainty in problems such as the Faint Young Sun Paradox \\citep[see e.g.,][]{Sagan1972, Goldblatt2011}. Earth, Venus, and to a small extent, Mars all have clouds, therefore we expect that exoplanets in the Habitable Zone (HZ) also harbor clouds in their atmospheres.  The climatic effects of clouds can be mimicked by increasing the surface albedo \\citep[see][and subsequent work]{Kasting1984}. Although this method can capture the climatic effects of clouds and reproduce the observed pressure-temperature structure of Earth, the measured global mean annual energy budget of Earth cannot be reproduced \\citep{Goldblatt2011}. Furthermore, the effects of clouds on the observed spectra is neglected by this approach. To address this issue, the measured cloud properties of Earth can be used in models assuming that the cloud properties do not change for different exoplanets \\citep{Kitzmann2010}. An alternative approach is to take the cloud heights and droplet sizes from observations, and fit the cloud fractions and liquid water paths in each cloud layer to match the global energy budget (GEB) of Earth \\citep{Goldblatt2011}. Such models adopt multiple cloud layers in agreement with observations \\citep{Rossow2005, Warren2007}. A parameterized cloud model is difficult to adopt for exoplanets which are not the exact duplicate of Earth because the number of cloud layers, and other cloud properties (such as the condensing material, height of the layers, droplet sizes, cloud fraction) are all unconstrained parameters in the general case.  Due to the large amount of observational and laboratory data available for Earth, detailed theoretical cloud models have been developed by the meteorological community and are available for 1D vertical \\citep[see e.g.,][]{Srivastava1967}, 2D \\citep[see e.g.,][]{Fan2007, Shima2009}, and 3D general circulation models \\citep[see e.g.,][]{Stephens2005}. We do not use these models because they are either too Earth-centric and study one specific cloud type on a local scale (e.g., cumulus or cirrus clouds only), and/or they are too computationally expensive and thus not suitable for parameter space exploration. Instead we develop a new 1D model using basic and well-understood physics also employed in these models.  We develop a new 1D microphysical cloud model which is suitable for a wide range of planetary and atmospheric parameters for exoplanets in the HZ. Our cloud model has five free parameters, and does not add a significant computational overhead to existing 1D atmosphere models.  The structure of the paper is as follows. We describe the numerical methods and codes used for the atmosphere modeling in Sect. \\ref{sec:nummet}. The optical properties of liquid and solid water droplets and our microphysical cloud model is discussed in Sect. \\ref{sec:cloudm}. We give the initial conditions of the simulations and perform convergence tests in Sect. \\ref{sec:tests}. We perform four sets of simulations to study the climatic effects of the five free parameters of the cloud model in Sect. \\ref{sec:res}. Finally, we discuss and summarize our results in Sects. \\ref{sec:disc} and \\ref{sec:sum}. ",
        "conclusions": "\\label{sec:disc} \\subsection{Critical view on the cloud model} \\label{sec:crit}  Often convection is treated as a 1D diffusion process in cloud models for brown dwarfs and giant planets \\citep[see e.g.,][]{Ackerman2001}. The formed cloud droplets can be displaced to both directions and this process artificially increases the vertical extent of the cloud layer. On Earth, we see that the regions where convection is upward and downward are physically separated, thus the downward motion does not influence the cloud properties. Therefore we adopted this concept in our microphysical model.  Cirrus clouds are produced by overrunning and undercutting on Earth (Sect. \\ref{sec:basic}). However, such horizontal motions cannot be self-consistently modeled in 1D codes, thus we are left to parameterize the height of the cloud deck. The adopted parametrization is probably the largest uncertainty in our model because it is solely based on Earth observations. The temperature or super-saturation level where cirrus clouds form could change, if the relative humidity, the circulation pattern of the atmosphere, etc changes. Unfortunately we are not aware of any cirrus model which could take such effects into account in a 1D model.  We consider cloud formation only by condensation, and consider one representative droplet size at a given height because the optical properties of a droplet size distribution can be approximated by a single droplet size \\citep{Hu1993}. We do not simulate coalescence and the vertical motion of a droplet population. The effects of such microphysical processes are included in a simplified way in the precipitation efficiency. This parameter could in principle be constrained self-consistently by simulating the vertical motion, condensation, and coalescence of a droplet population. However, such calculation would introduce more free parameters (e.g., the distribution and material of aerosols) and add a significant computational overhead to our numerical methods. To keep the free parameters limited and the computational time short, we introduce the precipitation efficiency.  The number density of aerosols is determined by various processes such as bubble bursts on the sea, wind lifting up dust from the continents, volcanos, etc. (see Sect. \\ref{sec:basic}). The total number of cloudy layers and the cloud fraction per cloud layer are determined by the circulation of the atmosphere. For example the rotation period, the friction between the surface and air, topography, relative humidity, number density of aerosols all influence the cloud fraction. The relative humidity is mostly determined by the topography of the planet, by the zonally averaged land/sea distribution and surface temperature profile. Detailed two dimensional cloud resolving models show that the precipitation efficiency depends on the water vapor content of the atmosphere, surface evaporation rate, amount of precipitating particles, zonal and vertical wind speeds, etc. \\citep{Sui2007, Gao2011}. As we either do not know the above described planet properties or cannot simulate the physical processes for exoplanets, we varied these parameters independently in Sect. \\ref{sec:res}. GCMs are able to resolve some of the correlations mentioned above, but such models are currently too computationally expensive to be used for parameter exploration. Therefore we were left to treat these parameters independently for now.  We use the critical Reynolds number in both cloud layers to limit the width of the cloud layers. A value of 200 is chosen for both layers, but other values could also be adopted. The critical Reynolds number determines the width of the cloud layer. If $Re$ is smaller(larger) than 200, the cloud layers become thiner(thicker) and a smaller(larger) value of precipitation efficiency would also reproduce the global average temperature and albedo of Earth. However, the critical Reynolds number should not be smaller than 70, because the flow around the droplets is laminar  and coalescence is not initiated in that case \\citep{Rossow1978}. Adopting too large values of $Re$ makes the cloud too thick geometrically and optically, and a precipitation efficiency of almost unity is necessary to reproduce the observed properties of Earth. We verified by simulations that the results are not sensitive to the exact value of $Re$. Varying the other cloud parameters yield to large trends, and the effect of different $Re$ is negligible compared to these trends.  We do not model the ice albedo feedback. The surface albedo decreases for higher surface temperature as the size of ice-covered areas become smaller due to melting. If the surface temperature decreases, the size of ice covered areas would increase, which in turn would further decrease the surface temperature. This is a positive feedback, as some cooling(warming) promotes further cooling(warming) of the surface. However, we keep the surface albedo fixed in all runs.    All our model runs were done with a fixed composition of non-condensable gases and a surface pressure of 1 bar. Comparison with Venus, Mars, and Titan indicates that the composition and mass of terrestrial planet atmospheres varies widely. Similar is expected for exoplanets. For Earth, the composition \\citep{Sagan1972, Owen1979} and possibly the mass \\citep{Goldblatt2009} of the atmosphere have evolved with the geochemical evolution of the planet. There are geochemical feedbacks between temperature and composition. Most famous is the temperature dependence of CO$_2$ concentration by silicate weathering, which is thought to exert a negative feedback on the global temperature on $10^5$ yrs timescales \\citep{Walker1981}. Until we have high resolution spectra for terrestrial exoplanets, the atmospheric composition will be difficult to constrain -- and modeling such effects is thus beyond the scope of our climate model.  \\subsection{Comparison of microphysical model to parameterized cloud models} \\label{sec:disc_cf} Our standard simulation is described in Sect. \\ref{sec:preceff} with a precipitation efficiency of 0.8. We compare this simulation to the standard simulations of \\cite{Goldblatt2011} and \\cite{Kitzmann2010}. Both of these cloud models are parameterized and have been applied to Earth as a test case.  \\cite{Kitzmann2010} uses two cloud layers, like our model. Their model reproduces the global average temperature of Earth, but the planetary albedo is somewhat smaller (0.27) than the observed value of 0.3. The planetary albedo is calculated as the ratio of incoming to outgoing solar flux in the GEB. This implies that the outgoing solar flux of \\cite{Kitzmann2010} is $10\\%$ smaller than the observed value. \\cite{Kitzmann2010} adopts 52 vertical grid cells and their cloud layer does not entirely fill one grid cell. The thickness of the grid cell containing the low level water cloud is for example 0.9 km, however the thickness of the cloud layer is 0.15 km (personal communication). \\cite{Kitzmann2010} adopts a droplet size distribution, number density, and cloud optical depth in accordance to observations. The cloud thickness is then derived from these quantities. The grid cell containing the cloud is determined by the observed pressure at the cloud top.  The cloud heights (or pressures) of \\cite{Goldblatt2011} are derived from observations. They adopt 30 vertical layers mostly centered around the troposphere and use three clouds layers. The low level liquid water cloud is resolved by 3 grid cells, the middle layer by 2, and the cirrus cloud layer by 1 grid cell. Their GEB values are shown in Fig. 5 of their paper and are in good agreement with the observed values. The largest relative error is $5\\%$ in their GEB fluxes with a typical error of 2-3$\\%$.  \\cite{Kitzmann2010} and our work adopt an initial droplet number density in agreement with observations. Our droplet number density is then reduced due to precipitation. The number density of droplets in \\cite{Goldblatt2011} is calculated using the liquid water path, average droplet size, the height of the cloud deck and top. Their lower liquid water cloud has for example 7 droplets per cm$^3$, which is lower than in the other two models. This difference is explained by the large vertical extent of their clouds (1.4 km, compared to 0.4 km in our standard model and 0.15 km in \\cite{Kitzmann2010}).  Finally, the results of parameterized cloud models differ in details from the results of the microphysical cloud model but all three models reproduce the GEB of Earth much better than clear sky models with an increased surface albedo. The advantage of our microphysical cloud model over parameterized models is that the cloud properties are determined self-consistently (except for the cirrus height), thus it is applicable to a much wider range of atmospheric and planetary properties than a parameterized model, which is tuned to Earth. The cloud model predict how cloud properties change and influence the climate, if relative humidity, cloud fraction, aerosol number density, atmospheric composition, etc is different than what is observed on Earth. Parameterized models have difficulty to cope with such issues.  \\subsection{Sensitivity of the climate} The simulations performed in Sect. \\ref{sec:res} map how the structure of the atmosphere is influenced by cloud fraction, aerosol number density, relative humidity and precipitation efficiency. For precipitation efficiency between 0 and 100\\%, the albedo and surface temperature varies between 0.47-0.15, and 259 K - 230 K, respectively. For a liquid water cloud fraction between 0 and 100\\%, the albedo and surface temperature varies between 0.16 - 0.58, and 305 K - 242 K, respectively. For an ice cloud fraction of 0 and 100\\%, the albedo and surface temperature varies between 0.26 - 0.3, and 282 K - 300 K, respectively. For aerosol number density between 70 cm$^{3}$ and 300 cm$^{-3}$, the albedo and surface temperature changes between 0.25 - 0.47, and 293 K - 262 K, respectively. The climate is least sensitive to the relative humidity. For relative humidity between 20\\% and 100\\%, the albedo and surface temperatures are 0.3 - 0.27, and 287 K - 294 K, respectively. A cloud-free atmosphere with zero relative humidity has a surface temperature of 273 K, and albedo of 0.17.  However, the relative humidity, cloud fraction and precipitation efficiency are not independent parameters. The relative humidity determines the amount of water vapor available for condensation and cloud formation. Therefore, assuming a zero cloud fraction and 77\\% relative humidity in Sect. \\ref{sec:cf} is probably unrealistic, if CCNs exist on the planet. Assuming less than 20\\% relative humidity and 40\\% cloud fraction in Sect. \\ref{sec:relhum}, and a precipitation efficiency of both zero and unity in Sect. \\ref{sec:preceff} are unlikely. The relative humidity and cloud fraction are dependent parameters, however the dependency is difficult to constrain even with global circulation models. Therefore we treat them as independent parameters. Determining the precipitation efficiency self-consistently would add new free parameters and significant computational overhead to the calculation (see Sect. \\ref{sec:crit}).  It shows the intricacy of the field that the albedo and climate depends rather strongly on the aerosols. It is also observed on Earth that the cloud properties are different above land, where the aerosol number density is high, and above sea, where the aerosol number density is low \\citep{Rossow1999}. Models to calculate the aerosol number density exists for Earth \\citep{Seinfeld2006} and it might be beneficial in the future to generalize those models for exoplanets."
    },
    "1208/1208.1520_arXiv.txt": {
        "abstract": "We report a measurement of the proton-air cross-section for particle production at the center-of-mass energy per nucleon of 57\\,\\TeV. This is derived from the distribution of the depths of shower maxima observed with the Pierre Auger Observatory: systematic uncertainties are studied in detail.  Analysing the tail of the distribution of the shower maxima, a proton-air cross-section of $\\left[505\\;\\pm22(\\text{stat})\\;^{+28}_{-36}(\\text{sys})\\right]\\,$mb is found. ",
        "introduction": "We present the first analysis of the proton-air cross-section based on measurements made at the Pierre Auger Observatory~\\cite{Abraham:2004dt}.  For this purpose we analyse the shape of the distribution of the largest values of the depth of shower maximum, $X_{\\rm max}$, the position at which an air shower deposits the maximum energy per unit of mass of atmosphere traversed.  The \\emph{tail} of the $X_{\\rm max}$-distribution is sensitive to the proton-air cross-section, a fact first exploited in the pioneering work of the Fly's Eye Collaboration~\\cite{Ellsworth:1982kv}.  To obtain accurate measurements of $X_{\\rm max}$, timing data from the fluorescence telescopes is combined with that from the surface detector array for a precise hybrid reconstruction of the geometry of events~\\cite{Sommers:1995dm}.  We place particular emphasis on studying systematic uncertainties in the cross-section analysis.  The unknown mass composition of cosmic-rays~\\cite{PedroICRC} is identified to be \\textit{the} major source of systematic uncertainty and accordingly the analysis has been optimised to minimise the impact of particles other than protons in the primary beam. This begins with restricting the analysis to the energy interval $10^{18}$ to $10^{18.5}\\,$\\eV, where the shape of the $X_{\\rm max}$-distribution is compatible with there being a substantial fraction of protons; also there are a large number of events recorded in this energy range. The corresponding average center-of-mass energy of a proton interacting with a nucleon is \\unit[57]{\\TeV}, significantly above the reach of the Large Hadron Collider. ",
        "conclusions": "We have presented the first measurement of the cross-section for the production of particles in proton-air collisions from data collected at the Pierre Auger Observatory.  We have studied in detail the effects of assumptions on the primary cosmic-ray mass composition, hadronic interaction models, simulation settings and the fiducial volume limits of the telescopes on the final result.  By analysing only the most deeply penetrating events we selected a data sample enriched in protons. The results are presented assuming a maximum contamination of 25\\,\\% of helium in the light cosmic-ray mass component. The lack of knowledge of the helium component is the largest source of systematic uncertainty.  However, for helium fractions up to 25\\,\\% the induced bias remains below $6\\,$\\%.  \\begin{figure}[tb] \\includegraphics[width=\\linewidth]{protonProtonInel.eps} \\caption{\\label{fig:pp} Comparison of derived $\\sigma_{pp}^{\\rm inel}$ to model predictions and {accelerator data~\\cite{lhc}}. Here we also show the cross-sections of two typical high-energy models, Pythia6~\\cite{Schuler:1993wr} and Phojet\\cite{Engel:1994vs}.  The inner error bars are statistical, while the outer include systematic uncertainties.} \\end{figure}  To derive a value of $\\sigma^{\\rm prod}_{p\\kern 0.1em\\text{-air}}$ from the measured $\\Lambda_\\eta$ we assume a smooth extrapolation of hadronic cross-sections from accelerator measurements to the energy of the analysis. This is achieved by modifying the model-predictions of hadronic cross-sections above energies of $10^{15}\\,\\eV$ during the air-shower simulation process in a self-consistent approach.   We convert the proton-air production cross-section into the total, and the inelastic, proton-proton cross-section using a Glauber calculation {that includes intermediate inelastic screening corrections}. In this calculation we use the correlation between the elastic slope parameter and the proton-proton cross-sections taken from the interaction models as a constraint.  We find that the inelastic proton-proton cross-section depends less on the elastic slope parameter than does the total proton-proton cross-section, and thus the systematic uncertainty of the Glauber calculation for the inelastic result is smaller. % {The data agree with an extrapolation from LHC~\\cite{lhc} energies to 57\\,TeV for a limited set of models.}    \\noindent\\textit"
    },
    "1208/1208.1716_arXiv.txt": {
        "abstract": "{X-ray emission from Active Galactic Nuclei (AGN) is dominated by the accretion disk around a supermassive black hole. The radio luminosity, however, has not such a clear origin except in the most powerful sources where jets are evident. The origin (and even the very existence) of the local bi-modal distribution in radio-loudness is also a debated issue.} {By analysing X-ray, optical and radio properties of a large sample of type 1 AGN and quasars (QSOs) up to $z> 2$, where the bulk of this population resides, we aim to explore the interplay between radio and X-ray emission in AGN, in order to further our knowledge on the origin of radio emission, and its relation to accretion.} {We analyse a large ($\\sim 800$ sources) sample of type 1 AGN and QSOs selected from the 2XMMi \\xmm\\ X-ray source catalogue, cross-correlated with the \\sdss\\ DR7 spectroscopic catalogue, covering a redshift range from $z\\sim 0.3$ to $z\\sim 2.3$. Supermassive black hole masses are estimated from the \\mgii\\ emission line, bolometric luminosities from the X-ray data, and radio emission or upper limits from the FIRST catalogue.} {Most of the sources accrete close to the Eddington limit and the distribution in radio-loudness does not appear to have a bi-modal behaviour. We confirm that radio-loud AGN are also X-ray loud, with an X-ray--to--optical ratio up to twice that of radio-quiet objects, even excluding the most extreme strongly jetted sources. By analysing complementary radio-selected control samples, we find evidence that these conclusions are not an effect of the X-ray selection, but are likely a property of the dominant QSO population.} {Our findings are best interpreted in a context where radio emission in AGN, with the exception of a minority of beamed sources, arises from very close to the accretion disk and is therefore heavily linked to X-ray emission. We also speculate that the radio-loud/radio-quiet dichotomy might either be an evolutionary effect that developed well after the QSO peak epoch, or an effect of incompleteness in small samples.} ",
        "introduction": "\\label{sect:intro}  Strong X-ray and radio emissions are properties that distinguish active galactic nuclei (AGN) from the whole population of galaxies. X-rays are the most direct manifestation of the accretion disk around a supermassive black hole (SMBH) at the centre of the galaxy hosting the AGN. Although radio emission is most apparent in a fraction of AGN, in particular in those classified as ``radio-loud'' (RL), which constitute $10-20$\\% of the local AGN population, recent work shows that even radio-quiet (RQ) AGN exhibit a radio-emitting core, which might result from some sort of radio plasma arising in the vicinity of the SMBH. While optical emission in AGN is due to the superposition of thermal components coming from different distances from the nucleus, with a contribution of radiation reprocessed outside the AGN central engine, both X-rays and radio emission can be used to probe the immediate environment of the SMBH.  From the observational point of view, two quantities appear most relevant in exploring the possible link between radio emission and accretion: the UV-based radio-loudness, defined as $\\mathcal{R}\\equiv\\pedix{F}{5\\,GHz,\\,rf}/\\pedix{F}{2500\\,\\AA,\\,rf}$ \\citep[monochromatic fluxes in the rest frame;][]{stocke92}, and the Eddington ratio $\\pedix{\\lambda}{Edd}\\equiv\\pedix{L}{bol}/\\pedix{L}{Edd}$, where $\\pedix{L}{bol}$ is the bolometric luminosity and $\\pedix{L}{Edd}\\equiv1.3\\times 10^{38}\\,\\pedix{M}{BH}/\\pedix{M}{\\sun}\\,$[\\lum] is the limiting luminosity of Eddington \\citep{eddington13,rees84}. Obtaining these quantities requires radio, UV and X-ray fluxes, plus reliable bolometric corrections and SMBH mass (\\pedix{M}{BH}) estimates. In contradiction to earlier works, \\citet{ho02} showed that $\\mathcal{R}$ is uncorrelated with \\pedix{M}{BH}, but strongly anticorrelated  with $\\pedix{\\lambda}{Edd}$. This was interpreted by \\citet{ho02} in the framework of changes of the accretion mode, from a radiatively efficient standard accretion at $\\pedix{\\lambda}{Edd}> 0.01$ to a radiatively inefficient ADAF (Advection Dominated Accretion Flow, which is prone to radio-emission) at lower Eddington ratios.  The local low-luminosity AGN population appears to show a dichotomy in radio-loudness, where two distinct populations appear to peak at $\\mathcal{R}\\sim 10-100$ (RL) and $\\mathcal{R}\\sim 0.1-1$ \\citep[RQ;][]{kellerman89}. This bimodality is not so apparent (or plainly non-existing) in deeper radio \\citep{white00} or X-ray selected surveys \\citep{brinkmann00}, where AGN samples display a continuous distribution in radio-loudness. It is then unclear whether there is something fundamentally different between strongly and weakly emitting radio AGN, and this is directly linked to the origin of radio emission in these objects.  With mass accretion rate largely regulating the luminosity of the AGN, the only other relevant parameter to modulate radio-loudness and whether an AGN develops or not a powerful jet, is the SMBH spin \\citep{blandford90,wilson95}. In a simple scenario where jets are formed in rapidly spinning SMBH, evolution by galaxy and associated SMBH mergers would naturally lead towards a SMBH population with low spin \\citep{berti08}. This would imply that smooth accretion (which tends to spin up SMBH) would be unimportant in the history of SMBH growth. However, the dependence of RL fraction on redshift and luminosity is still a strongly debated issue \\citep[\\eg,][and references therein] {jiang07}: the number of RL sources  in the analyzed samples, not large enough to study their two-dimensional distribution in redshift and luminosity, and the wide range of selection criteria used to define the samples observed contribute to a large range of contradicting results.  \\citet{sikora07} extend this view by showing that the anticorrelation between $\\mathcal{R}$ and $\\pedix{\\lambda}{Edd}$ comes in two parallel tracks, one for RL AGN residing in elliptical galaxies and one (lower) for RQ AGN residing in spiral galaxies. They propose a revised spin paradigm, where elliptical galaxies (and thence RL AGN) host highly spinning SMBH as a result of at least one major merger in the past, while spiral galaxies (and thence RQ AGN) underwent mostly chaotic accretion, \\ie, accretion of small mass fragments with random angular momenta. But since highly accreting luminous QSOs residing in ellipticals are largely RQ, speculations that radio emission might be intermittent have been put forward. Most recent versions of the spin paradigm call for retrograde systems (where the SMBH and the accretion disk counter rotate) as the mechanism to extract the most powerful jets \\citep{garofalo10}. RL AGN are mostly assumed to be retrograde, and RQ prograde. Natural evolution tends to make all SMBH-accretion disk systems prograde, which would explain the overwhelmingly large fraction of RQ AGN in the local Universe.  The release of large catalogues of fairly deep X-ray and radio sources, along with the optical photometry and spectroscopy provided by the Sloan Digital Sky Surveys \\citep[\\sdss; see][for the final public data release from \\sdss-II]{abazajian09}, calls now for a study of the relation between accretion and radio properties in large samples of AGN. This is of particular interest to infer the physical origin of the radio emission. In this study we use the incremental Second \\xmm\\ Serendipitous Source Catalogue \\citep[2XMMi;][]{watson09}, which we correlate with the \\sdss\\ Data Release~7 (DR7) to select a sample of X-ray detected type~1 AGN and QSOs. In this way, we use independent data to estimate the bolometric luminosity (from the X-ray data and suitable bolometric corrections) and the SMBH mass that we estimate from the \\sdss\\ spectra using the \\mgii\\ broad emission line. The latter effectively limits our sample to $z\\leq 2.3$, which is however enough to encompass the peak of the QSO activity epoch at around $z\\sim 2$. Note that we do not include type~2 AGN in this study, in part because the mass estimates based on other proxies (\\eg, the \\oiii\\ line width) are more uncertain and would compromise the homogeneity of our study. Radio information is obtained from the FIRST-VLA catalogue \\citep{first}.  The paper is organized as follows. In Section~\\ref{sect:sample} we describe the selection method used to generate the samples; Section~\\ref{sect:nucl} is devoted to recover the nuclear properties (\\ie, SMBH masses, nuclear bolometric luminosities, and Eddington ratios). In Section~\\ref{sect:rl} radio-loudness properties are discussed and compared with results from the literature. Finally, in Section~\\ref{sect:concl} we summarize our work. Throughout this paper, a concordance cosmology with $\\pedix{H}{0}=71\\,$km s$^{-1}$ Mpc$^{-1}$, $\\pedix{\\Omega}{\\ensuremath{\\Lambda}}=0.73$, and $\\pedix{\\Omega}{m}=0.27$ \\citep{spergel03,spergel07} is adopted. The energy spectral index, $\\alpha$, is defined such that $\\pedix{F}{\\ensuremath{\\nu}} \\propto \\apix{\\nu}{-\\ensuremath{\\alpha}}$. The photon index, defined as $\\pedix{N}{\\ensuremath{\\epsilon}} \\propto \\apix{\\epsilon}{-\\ensuremath{\\Gamma}}$, where $\\epsilon$ is the photon energy, is $\\Gamma=\\alpha+1$. ",
        "conclusions": "\\label{sect:concl}  In this paper, we explored the interplay between X-ray and radio emission in type~1 AGN, in order to investigate the origin of radio emission in the framework of the AGN Unification Model, and its relation with the different physical components of the central system as well as with the different accretion regimes in act. The availability of deep catalogues at different energies, from radio up to X-rays bands, with wide sky coverage, allow us to collect multiwavelength information for a large sample of $\\sim 800$ type~1 AGN, spanning a redshift range from $0.3$ to $2.3$.  For all the sources, we obtained the masses of the central black hole from the optical spectra, using the well-known relation between mass, emission-line width, and continuum luminosity. X-ray data were used to compute the bolometric output of the sources; being produced in the innermost regions of the central engine, the high-energy emission is one of the best proxies to estimate the bolometric luminosity, less affected by effects of reprocessing and external contamination than radiation emitted at larger distances. We tested our X-ray-based estimate of nuclear properties against the possible absorption in the X-ray band, finding that its effects would be negligible. In the previous sections, we discussed extensively the importance of the only contaminant we expect at high energy, \\ie\\ the emission from the jet in the most powerful radio sources.  Combining SMBH masses and bolometric luminosities, we recovered the Eddington ratios; the collection of radio information allow us to characterize the sample in terms of importance of the radio emission in the global energetic output. We note that one of the main characteristics of this work is the derivation of the different physical quantities from observations in different energy ranges, compared with the unavoidable dependence expected when observations in the same energy band are used (\\eg, the optical emission adopted both in the determination of the SMBH mass and as a proxy of the bolometric luminosity).  To assess whether these conclusions are a property related to the X-ray selection character of our sample, we tested our conclusions against a sample of FIRST radio AGN which appear in the \\sdss\\ catalogue, and for which X-ray information is obtained from the \\rosat\\ All-Sky Survey \\citep[RASS;][]{rass}.  Below we summarize our main results: \\begin{enumerate} \\item Our sources have typically $\\pedix{\\lambda}{Edd}> 0.01$. The sample analysed, which is effectively X-ray selected, might be biased towards high accretion rates. It is not surprising then that the trend of $\\mathcal{R}$ increasing towards decreasing $\\pedix{\\lambda}{Edd}$ noted in local samples is absent in our study, as we do not expect ADAFs or other radio-prone low-efficiency accretion modes to be present. We also find a few extreme cases (both in terms of $\\mathcal{R}$ and $\\pedix{\\lambda}{Edd}$), that we identify with beamed sources. \\item At a variance with lower-luminosity samples, ours does not show any hint of bimodality in radio-loudness. Despite the fact that our sample spans a wide redshift range, we have not found compelling evidence that bimodality develops with cosmic time. Although this cannot be excluded, we rather believe that the absence of a gap between RL and RQ is mainly due to the size of our sample, highly increased with respect to previous works. Our analysis suggests that, if the RL/RQ bimodality exist, it is a local effect. \\item We have computed X-ray loudness \\pedix{\\mathcal{R}}{X} for RL and RQ AGN (excluding strongly jetted sources) and we conclude that in the bulk of the AGN population, radio emission is tightly linked to the accretion disk and not to larger scale phenomena. \\end{enumerate}"
    },
    "1208/1208.3713_arXiv.txt": {
        "abstract": "Modeling of the self-consistent formation and evolution of disks as a result of prestellar core collapse reveals an intense early phase of recurrent gravitational instability and clump formation. These clumps generally migrate inward due to gravitational interaction with trailing spiral arms, and can be absorbed into the central object. However, in situations of multiple clump formation, gravitational scattering of clumps can result in the ejection of a low mass clump. These clumps can then give rise to free-floating low mass stars, brown dwarfs, or even giant planets. Detailed modeling of this process in the context of present-day star formation reveals that these clumps start out essentially as Larson first cores and grow subsequently by accretion. In the context of Pop III star formation, preliminary indications are that the disk clumps may also be of low mass. This mechanism of clump formation and possible ejection provides a channel for the formation of low mass objects in the first generation of stars. ",
        "introduction": "Numerical simulations that can track the evolution all the way from prestellar core collapse through the self-consistent formation of a disk and its subsequent long-term evolution are revealing new insights into star, brown dwarf, and planet formation. Our group has published a series of papers \\citep[e.g.,][]{vor06,vor10,bas12} that illustrate an early phase of disk evolution that is characterized by recurrent gravitational instability, and accretion driven by gravitational torques. Gravitational instability is triggered when accretion onto the disk drives the Toomre parameter of the disk below the critical value, resulting in the formation of gas clumps within nonlinear spiral arms. These clumps are generally driven inward through gravitational torques resulting from their interaction with trailing spiral arms. Some clumps are also dispersed due to tidal effects. Clumps that plunge inward to the central object can be invoked to explain the (FU Ori) luminosity bursts that are associated with young stellar objects \\citep{ken90}. Recent modeling \\citep{bas12} shows that the collapse of cores with sufficient mass and/or angular momentum can lead to disks with multiple clump formation in which the gravitational scattering of clumps leads to the ejection of low mass clumps. These clumps generally straddle the substellar limit, and may be the precursors of free-floating low mass stars, brown dwarfs, or even giant planets.  The above modeling is done in the thin-disk approximation, with a central sink cell of about 6 AU in radius, and a logarithmically spaced radial grid. These simplifications allow the modeling of the long-term evolution ( > 1 Myr) of the disk and core envelope structure. The combination of these features is generally not both feasible in fully three-dimensional simulations, even with adaptive mesh refinement. However, new three-dimensional simulations that resolve a large dynamic range of length scales have recently also found some of the same features, e.g., gravitational instability and episodic accretion \\citep{mac11}. ",
        "conclusions": "The model presented here has important implications for the formation of free-floating low mass stellar or substellar objects in the Galaxy. But what are the implications for the first generation of stars?  The clumps in the models with solar metallicity start out as essentially Larson first cores, defined as objects of a Jeans mass at the density where optical depth unity is first reached. This means that their initial mass is about $0.01\\, M_{\\odot}$ and they can grow subsequently by accretion from the disk, usually achieving a mass of about ten times this value, or $\\sim 0.1\\, M_{\\odot}$. Figure 2 shows the temperature-density relations for both $Z=Z_{\\odot}$ and $Z=0$ based on calculations of \\citet{omu05}. While optical depth unity is reached at a number density of $n \\simeq 10^{11}$ cm$^{-3}$ for solar metallicity, it is reached at $n \\simeq 10^{16}$ cm$^{-3}$ for the $Z=0$ case \\citep{omu05}. This is also the location of a local steepening of the temperature-density relation. The Jeans mass at this density and corresponding temperature $\\simeq 2000$~K is $\\sim\\, {\\rm few}\\, \\times 0.01\\, M_{\\odot}$. Preliminary calculations \\citep{vor12,des12} do find the formation of disk fragments of mass $\\sim 1\\, M_{\\odot}$, which is roughly consistent with the present-day simulations where the clumps grow to about ten times the initial Jeans mass. The formation, growth, and possible ejection of low mass Pop III stars through disk fragmentation is thus a tantalizing possibility that can be clarified by future calculations.  \\begin{figure}[h!] \\includegraphics[width=0.8\\textwidth]{./figure1.eps} \\caption{Gas surface density distribution (g cm$^{-2}$, log units) in a reference model at several time instances after the formation of a central star. The box size is 20,000 AU on each side and represents nearly the full extent of the computational domain. Arrows identify a clump that is ejected from the system after a multi-body interaction within the centrifugal disk.} \\label{fig:ejection} \\end{figure}  \\begin{figure}[h!] \\includegraphics[width=0.8\\textwidth]{./figure2.eps} \\caption{Temperature-density evolution for present day ($Z=Z_{\\odot}$) and primordial ($Z=0$) metallicities, based on the detailed thermal balance calculations of \\citet{omu05}} \\label{fig:barotropic} \\end{figure}   \\begin{theacknowledgments} We thank the organizers of First Stars IV for their efforts in putting together a highly stimulating conference. EIV acknowledges support from the RFBR grants 10-02-00278 and 11-02-92601. SB was supported by an NSERC Discovery Grant. \\end{theacknowledgments}"
    },
    "1208/1208.3463_arXiv.txt": {
        "abstract": "{ Within the cosmological concordance model, Cold Dark Matter (CDM) subhalos form the building blocks which merge hierarchically to more massive galaxies. This concept requires that even nowadays massive galaxies are surrounded by numerous subhalos. Since intergalactic gas is accreted by massive galaxies, observable e.g. as high-velocity clouds (HVCs) around the Milky Way, with extremely low metallicities, these can be suggested to represent the baryonic content of primordial Dark Matter (DM) subhalos. Another possibility of their origin is that they stem from disrupted satellite galaxies, but in this case, these gas clouds move unaccompanied by a bound DM structure.   } {HVCs are mainly observed at significant distances from the galactic disk, while those closer to the disk have on average lower infall speed and larger but not yet solar element abundances. This can be caused by interactions with the hot halo gas: decelerated by ram pressure and metal-enriched by the more abundant gas in the galactic halo. Since HVCs are observed with long gas tails and with irregular substructures, numerical models of gas clouds passing through hot halo gas are performed aiming at exploring their structure and compare them with observations. If HVCs are engulfed by DM subhalos, their gas must leave the DM gravitational potential and reflect this in their dynamics. On the other hand, the evolution and survival of pure gas models must be tested to distinguish between DM-dominated and DM-free clouds and to allow conclusions on their origin. } {A series of high-resolution hydrodynamical simulations using the adaptive-mesh refinement code FLASH V3.2 are performed for typical HVC masses, distances, and infall velocities. } {The models demonstrate that purely baryonic HVCs with low masses are disrupted by ram-pressure stripping and Kelvin-Helmholtz instabilities, while more massive ones survive, losing their initially spherical shape and develop significant substructures including cometary elongations in the column density distribution (``head-tail structure\"). On the contrary, HVCs with DM subhalos survive with more than 90\\% of their gas mass still bound and spherically shaped, approaching the Galactic disk like bullets. In addition, we find that velocity gradients along the cometary head-tail structures does not necessarily offer a possibility to distinguish between DM-dominated and purely gaseous HVCs. Comparison of models with observations let us conclude that HVCs are not embedded in a DM substructure and do not trace the cosmological subhalo population.} {}  \\keywords {ISM: clouds - Galaxy: halo - Galaxy: evolution - Cosmology: dark matter } ",
        "introduction": "In his general investigation of a possible Galactic corona, \\citet{1956ApJ...124...20S} was the first who predicted the existence of high-velocity clouds (HVCs). A few years later the systematic search for neutral hydrogen clouds succeeded and E. Raimond observed the 21 cm line of neutral hydrogen at high velocities \\citep{1963Natur.200..155M}.  Although the possibilities in observations, the theoretical understanding of the interstellar medium (ISM), and the facilities for numerical simulations have increased drastically, yet half a century later, the formation and evolution of HVCs are still under controversal discussion. The theories on their origin do not only vary in small details but lead to completely different pictures of the observed clouds, with tremendous consequences for galaxy formation and evolution. They are approaching the Galactic disk with high velocities and, depending on their distance, have to pass the hot coronal gas of the Milky Way halo.   Their radial velocities lie above or around escape velocity of the Milky Way and their rotation component is incompatible with the Galactic rotation \\citep{1999ApJ...514..818B}. This obtrudes the extragalactic origin of these gas clouds, excluding the widely discussed possibility of  ejecta from galactic fountains \\citep{1976ApJ...205..762S, 1980ApJ...236..577B}.   This issue leaves open two further fundamental possibilities: They are: 1) relics of cosmological structure formation in a Cold Dark Matter (CDM) cosmology \\citep{2006ApJ...646L..53C} and therefore distant ``building blocks\" from galaxy formation in the Local Group \\citep{1999ApJ...514..818B} or 2) remnants from galaxy collisions, from tidal disruption of dwarf galaxies \\citep{2004ApJ...603L..77P}, or/and from close-by tidal interactions of satellite galaxies \\citep{1996MNRAS.278..191G, 1998MNRAS.299..611B}. The main difference to distinguish between both sources is their Dark Matter (DM) content: in case 1) clouds should be strongly dominated by DM, but not in 2). Decisive issues could be derived from detailed observations in connection with complex numerical simulations, accounting for their distances, physical state, dynamical structure and chemical abundances.   \\begin{table*}[htdp] \\caption{Distances and Abundances of HVC and IVC} \\begin{center} \\begin{tabular}{lllc} \\hline \\hline Cloud \t\t&\t\tDistance \t&\tAbundance\t\t& References \\\\ &\t\t(kpc)\t\t&\t(rel. to solar)\t&\t\\\\ \\hline Complex A\t&\t\t4.0 - 9.9 \t&\t0.02 - 0.4\t\t\t& 1\\\\ Complex C \t&\t\t$10 \\pm 2.5$\t&\t0.09 - 0.29\t& 2,3  \\\\ Complex H \t&\t\t$27 \\pm 9$&\t\t\t\t\t& 4 \\\\ CHVC 224.0-83.4-197&\t?\t\t&\t$< 0.46$\t\t\t& 5 \\\\ CHVC 125+41-207&\t?\t\t&\t$<0.2$\t\t\t& 6 \\\\ Upper IV Arch\t&\t\t0.8 - 1.8 \t&\t$\\approx 1$\t\t& 1\\\\ Lower IV Arch \t&\t\t0.4 - 1.9\t&\t$\\approx 1$ \t\t& 1\\\\ LL IV Arch \t&\t\t0.9 - 1.8 \t&\t$1.0 \\pm 0.5$ \t\t& 1\\\\ Complex K\t&\t\t$< 6.8$\t&\t$< 2.0$\t\t\t& 1\\\\ Pegasus-Pisces Arch &\t$< 1.1$\t&\t$0.54 \\pm 0.04$\t& 1\\\\ Complex gp\t&\t\t$0.8-4.3$\t&\t1-2\t\t\t\t& 1\\\\ \\hline \\end{tabular} \\tablebib{ (1)~\\citet{2001ApJS..136..463W}; (2)~\\citet{2008ApJ...684..364T}; (3)~\\citet{2007ApJ...657..271C}; (4)~\\citet{2003ApJ...591L..33L}; (5)~\\citet{2002ApJ...572..178S}; (6)~\\citet{2001A&A...370L..26B}. } \\end{center} \\label{tab:metall} \\end{table*}%   \\subsection{Observations}  \\subsubsection{Distances, All-Sky Surveys, General Properties}  Detailed studies measure distances of a few kpc for the extended HVC complexes \\citep{2006ApJ...638L..97T, 2008ApJ...672..298W}, e.g. Complex C with a distance of $d = 10\\pm2.5\\, \\mbox{kpc}$ \\citep{2008ApJ...684..364T}. For the massive Complex H a distance of 27 kpc \\citep{2003ApJ...591L..33L} and a mass of $2\\times 10^7\\,M_{\\odot}$ \\citep{2006ApJ...640..270S} is derived. Beside such large complexes which cover around 40\\% of the sky for a resolution limit of $N_{HI}\\ge 7\\times 10^{17}\\, \\mbox{cm}^{-2}$ \\citep{1995ApJ...447..642M}, there have also been found isolated, compact gas cloudlets with a small angular extent \\citep{1999A&A...341..437B}. The first all-sky study of this sub-population, called Compact High Velocity Clouds (CHVCs), was presented by \\citet{2002A&A...392..417D} as a combination of data extracted by \\citet{2002A&A...391..159D} from the Leiden Dwingeloo Survey for the northern hemisphere and from the HI Parkes All Sky Survey (HIPASS), extracted by \\citet{2002AJ....123..873P} for the southern hemisphere. Due to the compactness of CHVCs, the lack of distance indicators, and the poor knowledge about their physical environment, direct distance measurements are difficult.  Observations of M31 have shown recently that HVC-analogs are clustered around their host galaxy with a maximum distance between 50 kpc \\citep{2004ApJ...601L..39T} and 60 kpc \\citep{2005A&A...432..937W}. Most recently, \\citet{2011A&A...528A..12R} have studied Ca II absorbers at low redshift and propose that the characteristic radial extent of HVC analogs with $\\log N(HI)\\ge 17.4$ in galaxies with $z<0.5$ is $R_{HVC}\\approx 55\\,\\mbox{kpc}$.  Another remarkable property of HVCs is a 2-phase structure with a cold, dense core and a warmer, more tenuous envelope, which was already discussed by \\citet{1991A&A...250..484W} for clumps in the larger complexes and is still supported by more recent observations of smaller cloudlets as in \\citet{2006A&A...457..917B}. \\citet{2010ApJ...708L..22G} presented ALFALFA (Arecibo Legacy Fast ALFA) observations of ultra-compact high velocity clouds (UCHVCs) at assumed extragalactic distances. The derived masses of the measures UCHVCs at a distance of 1 Mpc are below $10^6\\, \\Msun$, which explains the non-detection by previous HI surveys, such as \\citet{2007ApJ...662..959P}.  \\subsubsection{Dark Matter Minihalo and Ram Pressure Stripping: \\\\}  If the CHVCs permeate the intergalactic space their origin allows two possibilities:  they are disrupted remnants from galaxy collisions and formed in the tidal tails or they represent the relics of structure formation in a CDM cosmology. The latter possibility would provide the requested link to the postulated DM building blocks of hierarchical galaxy formation which means that this baryonic cloud mass should be harbored by a so-called DM mini-halo. \\citet{2009ApJ...707.1642N} have investigated the orbit of the Smith Cloud under the consideration that the cloud is confined by a DM halo.  Detailed observations of HVCs are showing a head-tail structure in the hydrogen column density distribution in both over-densities within the large complexes \\citep{2000A&A...357..120B} and single CHVCs, like presented for example by \\citet{2001A&A...370L..26B} and \\citet{2006A&A...457..917B}. In agreement with theoretical and numerical studies this structural properties are unambiguously indicative for gas clouds moving through a surrounding plasma. It is expected that if the force due to ram pressure exceeds the gravitational restoring force that keeps the cloud material via self-gravity together, material is stripped off the cloud \\citep{1972ApJ...176....1G}. This leads to a head-tail structure of the cloud, which is observed with radio observations of the 21cm line of neutral hydrogen \\citep{2000A&A...357..120B,2002AJ....123..873P,2005ASPC..331..105W}. Furthermore, the mixing of hot halo and cool cloudy gas might be responsible for the observed ionization stages in the hot halo \\citep{1996ApJ...458L..29S}.  Observationally derived structure parameters allow the conclusion that CHVCs are bound by a DM halo of significant mass \\citep{2000A&A...354..853B}. While the gravitational potential in the vicinity of the HVC is dominated by the combined gravitational potentials of the cloud itself and the host galaxy, a DM halo around the CHVC enlarges the total mass of the cloud-halo system and therefore also the volume in which the cloud potential exceeds the potential of the major galaxy, so that HVC material is gravitationally stronger bound to the cloud.  \\begin{table*}[htdp] \\caption{Initial parameters of the high-velocity clouds in different simulations. First column: initial setup no. for every run; col. 2: relative HVC velocity $v_w$; col. 3-5: HVC temperature $T_{gas}$, radius $R_{gas}$ in pc, and mass $M_{gas}$ in $\\Msun$; DM subhalo scale radius of $r_0$ (col. 6) in pc and mass $M_{DM}$ (col. 7) in $\\Msun$; col. 8: baryon-to-total mass ratio $M_{gas}/M_{tot}$;  maximum refinement level of the adaptive grid (Max.Ref.) in col. 9 and effective resolution (Res.) in col. 10. In the last column (Run Nr.) a unique number is assigned to every simulation run.}  \\begin{center} \\begin{tabular}{lcccccccccr} \\hline \\hline Setup & $v_{w}$ & $T_{gas}$ & $R_{gas}$\t& $M_{gas}$ & $ r_0$ & $M_{DM}$\t & $M_{gas}/M_{tot}$ & Max.Ref. & Res. & Run\\\\  & [$km\\,s^{-1}$] & $[K] $  & $[pc]$ & $[\\Msun]$\t& $[pc]$ & $ [\\Msun]$ & {\\bf ($M_{tot}=M_{gas}+M_{DM}$) }  &  & [pc] & Nr.\\\\  \\hline  2\t& 200&\t1000\t&\t194\t&\t$8.55 \\times 10^5$\t&\t   -\t&\t-\t\t\t\t&\t-\t& 5\t& 38\t&I \\\\\t% & 200&\t1000\t&\t194\t&\t$8.55 \\times 10^5$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.15\t& 5\t& 38\t&II\\\\\t% \\hline 2-D\t& 200&\t1000\t&\t125\t&\t$4.12 \\times 10^5$\t&\t-\t& \t-\t\t\t\t&\t-\t& 5 \t& 38\t&III \\\\\t% & 200&\t1000\t&\t125\t&\t$4.12 \\times 10^5$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.08\t& 5\t& 38\t&IV\\\\\t% \\hline 2-DD& 200&\t1000\t&\t127\t&\t$4.40 \\times 10^5$\t&\t-\t&\t-\t\t\t\t&\t- \t& 6\t& 25\t&V\\\\\t% & 200&\t1000\t&\t127\t&\t$4.40 \\times 10^5$\t&\t-\t&\t-\t\t\t\t&\t- \t& 7\t& 5\t&VI\\\\\t% & 200&\t1000\t&\t127\t&\t$4.40 \\times 10^5$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.08 \t& 7\t& 5\t&VII\\\\ \t% & 200&\t1000\t&\t127\t&\t$4.40 \\times 10^5$\t&\t250\t&\t$1.63\\times 10^7$\t&\t0.03 \t& 6\t& 25\t&VIII\\\\\t% & 200&\t1000\t&\t127\t&\t$4.40 \\times 10^5$\t&\t250\t&\t$1.63\\times 10^7$\t&\t0.03 \t& 7\t& 5\t&IX\\\\\t% \\hline & 250&\t1000\t&\t127\t&\t$4.40 \\times 10^5$\t&\t-\t&\t-\t\t\t\t&\t- \t& 6\t& 18\t&X\\\\\t% & 250&\t1000\t&\t127\t&\t$4.40 \\times 10^5$\t&\t250\t&\t$1.63\\times 10^7$\t&\t0.03 \t& 6\t& 18\t&XI\\\\\t% \\hline 3\t& 200&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t-\t&\t-\t\t\t\t&\t- \t& 4 \t& 101&XII\\\\  \t% & 200&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t-\t&\t-\t\t\t\t&\t- \t& 5 \t& 50\t&XIII\\\\ \t% & 200&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t-\t&\t-\t\t\t\t&\t- \t& 6 \t& 25\t&XIV\\\\ \t% & 200&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t-\t&\t-\t\t\t\t&\t- \t& 7 \t& 13\t&XV\\\\ \t% & 200&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.39 \t& 5 \t& 50\t&XVI\\\\  \t% & 200&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.39 \t& 6 \t& 25\t&XVII\\\\  \t% & 250&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t-\t&\t-\t\t\t\t&\t- \t& 5 \t& 50\t&XVIII\\\\ \t% & 250&\t2000\t&\t376\t&\t$3.10 \\times 10^6$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.39 \t& 5 \t& 50\t&XIX\\\\ \t% \\hline 3-D\t& 200&\t2000\t&\t230\t&\t$1.39 \\times 10^6$\t&\t-\t&\t-\t\t\t\t&\t- \t& 5\t& 50\t&XX\\\\ \t% & 200&\t2000\t&\t230\t&\t$1.39 \\times 10^6$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.22 \t& 5\t& 50\t&XXI\\\\ \t% & 200&\t2000\t&\t230\t&\t$1.39 \\times 10^6$\t&\t150\t&\t$4.93\\times 10^6$\t&\t0.22 \t& 6\t& 25\t&XXII\\\\ \t% \\hline \\hline \\end{tabular} \\end{center} \\label{tab:overview} \\end{table*}%  In addition, whether HVCs represent the baryonic content of cosmological DM sub-halos or not, should become visible during the ram pressure stripping (RPS) action on their passage through the hot halo gas. The gravitational acceleration in the vicinity of the HVC is dominated by the self-gravity of the cloud and the external gravitational potential of its host galaxy. At first, the cloud gas is confined within the equipotential of cloud and galaxy potential. Once the stripped-off material runs over this region, it gains an additional acceleration in comparison to the material that is stripped off but still gravitationally bound to the cloud. This effect is expected to be hardly noticeable if the HVC's total mass is only determined by its gas mass, because the potential of the massive host galaxy mass would widely dominate the self-gravity of the cloud. If the cloud is embedded in a DM subhalo the total mass of the HVC and therefore its self-gravity is much higher and could make this effect visible in velocity-position space as a sharp change in the velocity gradient along the head-tail structure of the cloud, as observations obtrude \\citep{2000A&A...357..120B,2001A&A...370L..26B,2006A&A...457..917B}.     \\subsection{Numerical Simulations}  \\subsubsection{Ram Pressure Stripping: \\\\}  \\citet{2001ApJ...555L..95Q} have studied by means of hydrodynamical simulations the behavior of CHVCs passing through the dilute hot halo gas of a typical disk galaxy at large distances. The main purpose of their numerical models was to compare the emerging head-tail structure with observations, in order to draw conclusions on the necessity of a DM component and on a possible lower density limit for the surrounding intergalactic medium (IGM).  They found head-tail structures in environments with a density higher than $10^{-4}\\, cm^{-3}$, whose lifetimes in simulations without a DM halo are very short, $\\approx 10\\,$Myrs only, while the tails in their DM dominated clouds survive for $\\approx 1\\,$Gyr.  Although the authors have simulated a grid of HVC models, it is worth to point out, that there is no comparison of the same cloud with and without DM, like we discuss in this work, but the authors study two different scenarios: HVCs as close-by structures at a distance of 1-10 kpc and masses between 10 and 100 $M_{\\odot}$ vs. very distant DM dominated clouds at a distance of 300 kpc, a radius of 1 kpc and a mass of more than $10^7\\, M_{\\odot}$, contrary to the typical maximum distance of around 55 kpc \\citep{2011A&A...528A..12R}.  A snapshot of the DM dominated cloud modeled over 300 Myrs (Fig. 2 of \\citet{2001ApJ...555L..95Q}) with a constant velocity of 200 km/s travels over a distance of 60 kpc through the galactic halo, i.e. must have faced changes in the ambient densities, a tidal field and an acceleration by the host galaxy. Due to the lack of self-gravity, low-mass DM-free gas clouds are disrupted very quickly so that the evolution is only shown after 3 Myrs \\citep[][Fig. 3]{2001ApJ...555L..95Q}. The DM-free clouds in the simulation runs presented in this work can easily survive for 100 Myrs, not only because of the included self-gravity but also due to the much higher masses and distances, in agreement with recent observations.  \\citet{2009ApJ...698.1485H} have shown with grid-based hydrodynamic simulations that smaller clouds with HI masses $<10^{4.5}\\, \\Msun$ will lose their HI content during their paths of 10 kpc at maximum (for typical relative velocities and halo densities). They propose that the stripped material which is usually gas in the warm and ionized phase  can contribute to the extended layer of warm, ionized gas at lower galactic heights \\citep{1993AIPC..278..156R, 2008PASA...25..184G} and later on recool and form smaller HI cloudlets which are classified as low or intermediate velocity clouds (LVC, IVCs).    \\subsubsection{Heat Conduction: \\\\}  Due to the large temperature gradient between the cloud and the surrounding halo gas, saturated heat conduction should be an important process of energy transport between the gas phases, leading to the classical two-phase structure. Furthermore, heat conduction stabilizes clouds with head-tail structure and is extending their lifetime \\citep{2007A&A...472..141V}. In addition, \\citet{2007A&A...475..251V} have shown that in models which combine heating, cooling and saturated heat conduction with self-gravity, condensation of halo gas on the cloud surface can dominate evaporation and lead to a net accretion of ambient gas.  For the warm cosmic rain theory by \\citet{2009ApJ...698.1485H} this means that the stripped off material is not only warmer and higher ionized than the cold core of HVCs, but would have other significant consequences: On average HVCs contain much smaller element abundances than the Milky Way ISM (see Tab.~\\ref{tab:metall}) what clearly advocates for their primordial origin. HVCs affected by heat conduction and stripping in the presumably metal-rich halo gas are expected to have also a higher metallicity because the stripped cloud surface consists not only of cloud gas but also of accreted halo gas. If part of the LVC and IVCs are the decelerated survivors of disrupted or ablated HVCs they should have a higher metallicity than the less processed HVCs, which is observationally confirmed (Tab.~\\ref{tab:metall}). Since IVCs/LVCs rain down with velocities below escape velocity, an origin from the galactic fountain cannot be ruled out. The cores of HVCs, which survive ram-pressure stripping, can be decelerated by either the drag of the hot halo gas or other processes like the buffering of plane-parallel galactic magnetic field lines which have to couple to the HVC plasma \\citep{1997A&A...320..746Z}. This effect would contribute to the heating of the halo gas \\citep{2011CoPhC.182.1784J}.  In a recent study about the origin of the HVCs \\citet{2009MNRAS.397.1804B} suggested that they can be formed by thermal instability in the warm-hot Galactic halo medium. For Milky Way sized halos, however, they found that this happens only in regions outside of 100 kpc and for an almost perfectly flat entropy profile, when thermal instabilities are not damped by a combination of buoyancy and thermal conduction. They concluded that it is therefore a rather unlikely formation scenario for HVCs.   \\begin{table}[tdp] \\caption{ The initial distribution of the cloud follows in all setups a hydrostatic equilibrium, which is derived by Eq.~\\ref{equ:hydro2}, where $M(r)$ includes either only the gas mass or in addition to that also the enclosed DM mass following Eq.~\\ref{equ:burkert1} and \\ref{equ:burkert2} with a scale radius of $r_0 = 150\\, pc$ or $r_0 = 250\\, pc$. } \\begin{center} \\begin{tabular}{lll} \\hline \\hline Setup \t\t&\t\tT [K] \t&\t$M(r)$ \\\\ \\hline 2\t\t\t&\t\t1000 K\t\t&\tgas only\t\t\\\\ 2-D\t\t\t&\t\t1000 K\t\t&\tgas + DM halo ($r_0 = 150\\,pc$)\t\\\\ 2-DD\t\t&\t\t1000 K\t\t&\tgas + DM halo ($r_0 = 250\\,pc$)\t\\\\ 3\t\t\t&\t\t2000 K\t\t&\tgas only\t\t\\\\ 3-D\t\t\t&\t\t2000 K\t\t&\tgas + DM halo ($r_0 = 150\\,pc$)\t\\\\ \\hline \\end{tabular}  \\end{center} \\label{tab:setups} \\end{table}%  \\subsection{General Questions}  Since differences in the ram pressures stripped gas and in the head-tail structure are not yet explored with respect to their possible origins, here we present numerical models of HVCs with and without a stabilizing DM subhalo based on the following questions: Is it possible for the cloud with DM to keep the gas longer bound as without? Does a DM halo suppress hydrodynamical instabilities? Is a cloud without a DM halo already disrupted and does it lose most of its gas due to the ram pressure produced by its supersonic motion (up to more than $400\\kms$) through the hot gaseous halo of its host galaxy or in the outskirts of the dilute halo gas? Can we distinguish from the observed velocity - position distribution of the head-tail structure whether the observed cloud contains DM? If they are embedded in a DM halo, are HVCs the low mass extensions of gas rich subhalos with (ultra faint) dwarf galaxies as their higher mass analogs? Moreover, if they are DM-free and a significant fraction of the few observed satellite galaxies have a tidal origin \\citep{2006A&A...456..481B,2000ApJ...543..149O}: where are all the DM satellites we see in large cosmological simulations?  \\begin{figure}[htbp] \\begin{center} \\includegraphics[width=\\linewidth]{images/setups.eps} \\caption{Initial density (top panel) and pressure (lower panel) distribution of the model setups described in Tab.~\\ref{tab:overview}.} \\label{fig:overview} \\end{center} \\end{figure}  \\subsection{Paper Structure}  This paper has the following structuring: First the numerical method and the model setup of both the host galaxy and the HVCs are presented. In Sect.~\\ref{sec:lagrange} we discuss the size of the Roche-lobes in this 2-body system and the consequences for the gas that is lost by RPS.  The effect of RPS that leads to the observed cometary appearance of the clouds is discussed in Sect.~\\ref{sec:rampres}, while further possible mass-loss procedures are mentioned in Sect.~\\ref{sec:other}.  Because HVCs do not necessarily fall directly towards the observer nor radially into the center of the Milky Way we take a close look on the line-of-sight problem in Sect.~\\ref{sec:lineofsight}. Not only the column density distribution varies when changing the inclination between object and observer but the velocity measurements can cover only the radial velocity. Therefore the whole phase-space distribution of the observed clouds is highly sensitive to line-of-sight effects. For all performed simulations we processed the data in a way that we simulate different line of sight directions to look for signatures that are dependent on the line of sight.  Finally, the dependency of the results on parameters like the initial velocity (see Sect.~\\ref{sec:windvel}) and the DM mass (see Sect.~\\ref{sec:dmmass}) and shape (Sect.~\\ref{sec:dmshape}) are discussed.   We have included not only a self-gravitating gas cloud as well as a static gravitational potential that represents a DM mini-halo but also an external gravitational potential which gives us the opportunity to study the equipotential lines of the combined potential of the galaxy and the HVC.  \\begin{figure*}[tbp] \\begin {minipage}[b]{0.47\\linewidth} \\includegraphics[width=\\linewidth, bb = 0 0 566 255,clip]{images/equi_setup2-DD_without_000_dm_ref7_0001.eps} \\end {minipage} \\begin {minipage}[b]{0.47\\linewidth} \\includegraphics[width=\\linewidth, bb = 0 0 566 255,clip]{images/equi_setup2-DD_with_250_dm_ref7_0001.eps} \\end {minipage}  \\begin {minipage}[b]{0.47\\linewidth} \\includegraphics[width=\\linewidth, bb = 0 0 566 255,clip]{images/equi_setup2-DD_without_000_dm_ref7_0100.eps} \\end {minipage} \\begin {minipage}[b]{0.47\\linewidth} \\includegraphics[width=\\linewidth, bb = 0 0 566 255,clip]{images/equi_setup2-DD_with_250_dm_ref7_0100.eps} \\end {minipage} \\caption{Column density distribution of the gas cloud overlaid with equipotential lines for Setup 2-DD (see Tab.~\\ref{tab:overview} for its parameters) at a maximum refinement level of 7 (effective spatial resolution: 5 pc): without DM halo (left panels, Run VI) and with DM halo of scale radius $r_0=250\\,pc$  (right panels, Run IX). The contour lines of the greyscale \\HI column density distribution range between $\\mbox{log}(N_{\\HI}) = 18.8 \\mbox { ... } 21.2 \\, cm^{-2}$ in steps of 0.3 dex. Overlaid are the equipotential lines of the total gravitational potential by the gaseous cloud, the DM subhalo of the cloud (if present), and the external galaxy (see text) in steps of $2.5\\times 10^{11}\\,cm^{2}\\,g^{-2}$. In red, the equipotential line of the innermost Lagrangian point is depicted. Top panels show the initial distributions, bottom panels after a simulation time of 100 Myrs. (Step-like contour lines reflect the coarser grid further away from the cloud.)} \\label{fig:rochelobe2dd} \\end{figure*}   The initial distance between the HVC and the external mass was chosen to be 50 kpc, which leads to cloud masses in the order of $10^5$ to $10^6$ solar masses. Outside of 60 kpc the HVCs are either fully ionized and therefore not visible in neutral hydrogen or HVCs origin from an interaction scenario like the Magellanic Stream which would also explain the typical maximum distance of around 50 - 60 kpc. ",
        "conclusions": "Summarizing, the following major issues from the presented series of model runs can be derived. While it is possible to reproduce the observed change in the velocity gradient, this feature is not directly connected to the Roche lobe and therefore the position of the observed ``knee\" in velocity-position phase space cannot help to distinguish between DM dominated and pure gas clouds. For a variety of initial distributions we can create these phase space distributions by changing the inclination of the line-of-sight. In a further study, clouds with a significant azimuthal motion must be tested with respect to their velocity gradient of the stripped gas for which we expect a more significantly observable velocity ``knee\".  We simulate DM-embedded clouds with baryonic to total mass ratios between 0.03 and 0.39 and found that the results depend only marginally on the DM mass and profile and that already a DM halo with a mass of 1.6 times the baryonic mass stabilizes the gas cloud against ram-pressure disruption.  After 100 Myrs the column-density distributions of the clouds with DM halo can be well distinguished from pure gas clouds. While the isodensity lines of the core of gas clouds without an additional DM gravitational potential are losing their initial shape after half the simulation time and show a lot of substructures, for all DM-dominated models HVCs remain spherical and centrally concentrated, even if the DM content is low as in Setup 3.  Furthermore an extended head-tail structure at galactic heights between 30 and 50 kpc can only be reproduced for DM-free clouds.  In our simulations, clouds with DM - with baryonic masses in the range from a few times 10$^5$ to a few times $10^6 \\Msun$ - approach the galactic disk like bullets with high velocities, while DM-free clouds are already decelerated and distorted at higher galactic heights. Observed HVCs with distances up to 30 kpc, like e.g. Complexes C and H, reveal a very heterogeneous multi-phase internal gas structure without any central concentration, which we only obtain in simulations without DM.  Furthermore, the observational lack of HVCs close to the Milky Way disk (only LVCs exist) and also that high-speed infalling \\HI gas is not observed in galaxies with extended \\HI envelopes, as e.g. NGC~2403 \\citep{2002AJ....123.3124F} let us conclude that HVCs are decelerated to LVCs or disrupted on their infall into galactic halos, contrary to the simulated ``bullets\" of clouds which have an accompanying DM minihalo.  Our DM-free models can both explain the proximity of the large HVC complexes to the galactic disk even though they start their journey from a height of 50 kpc above the plane, as well as their morphologies, like e.g. from elongated shapes for more massive clouds to head-tail structures and the inherent amount of substructures.  Our conclusions contain the following further serious consequences for the cosmological model and for our understanding of cosmological structure formation.  \\begin{enumerate} \\item\t Since HVCs are DM-free, they cannot be the cosmological small-scale relics expected from $\\Lambda$CDM structure. \\item Since HVCs have typical distances to the disk of their host galaxy of not more than 50 - 60 kpc and do not populate the intergalactic space frequently enough, it is obtruding that, if they are DM-free, they originate from stripped-off gas of satellite galaxies or from condensations of diffuse intergalactic medium due to thermal instability rather than from an cosmological background.  \\end{enumerate}"
    },
    "1208/1208.4844_arXiv.txt": {
        "abstract": "We show that the position of the central dark matter density peak may be expected to differ from the dynamical center of the Galaxy by several hundred parsec. In Eris, a high resolution cosmological hydrodynamics simulation of a realistic Milky-Way-analog disk galaxy, this offset is \\mbox{300 - 400} pc ($\\sim 3$ gravitational softening lengths) after $z=1$. In its dissipationless dark-matter-only twin simulation ErisDark, as well as in the Via Lactea II and GHalo simulations, the offset remains below one softening length for most of its evolution. The growth of the DM offset coincides with a flattening of the central DM density profile in Eris inwards of $\\sim 1$ kpc, and the direction from the dynamical center to the point of maximum DM density is correlated with the orientation of the stellar bar, suggesting a bar-halo interaction as a possible explanation. A dark matter density offset of several hundred parsec greatly affects expectations of the dark matter annihilation signals from the Galactic Center. It may also support a dark matter annihilation interpretation of recent reports by \\citet{weniger_tentative_2012} and \\citet{su_strong_2012} of highly significant 130 GeV gamma-ray line emission from a region $1.5^\\circ$ ($\\sim200$ parsec projected) away from \\mbox{Sgr A*} in the Galactic plane. ",
        "introduction": "Dissipationless (dark matter only) N-body simulations predict a nearly universal density profile of dark matter (DM) halos, the so-called NFW profile \\citep{navarro_universal_1997}, which features a central $1/r$ density cusp. If the cooling and condensation of gas inside these halos \\citep{white_core_1978,fall_formation_1980} is gradual, then so-called ``adiabatic contraction'' \\citep{blumenthal_contraction_1986,gnedin_response_2004} will pull DM into the central regions, thereby increasing the central DM density and further steepening the slope of its density profile. It is thus natural to expect the maximum of the DM density to occur at the dynamical center of a galaxy, the lowest point of its gravitational potential. This expectation has made the Galactic Center (GC) a preferred target for indirect DM detection efforts searching for an annihilation signal \\citep{bergstroem_observability_1998,gondolo_dark_1999,aharonian_hess_2006,baltz_pre-launch_2008,danninger_searches_2012}.  As data from the Fermi Gamma-ray Space Telescope \\citep{atwood_large_2009} has been accumulating, the number of studies reporting gamma-ray ``excesses'' or ``anomalies'' from the GC, that could be interpreted as a DM annihilation signal, has steadily grown \\citep[see e.g.][]{goodenough_possible_2009,hooper_dark_2011,hooper_origin_2011,abazajian_detection_2012}. These analyses have searched for the broad gamma-ray continuum signal thought to arise from the decay and hadronization of DM annihilation products \\citep{bergstroem_observability_1998}. While this is expected to be the dominant DM annihilation signature, it is unfortunately difficult to distinguish it from conventional sources (e.g. milli-second pulsars, Abazajian 2011\\nocite{abazajian_consistency_2011}, or cosmic-ray interactions with molecular clouds, Yusef-Zadeh et al. 2012\\nocite{yusef-zadeh_interacting_2012}), and at present a purely astrophysical explanation of these signals cannot be excluded.  More recently, however, there have been surprising reports of a highly statistically significant line-like feature at $\\sim 130$ GeV in Fermi data from the GC \\citep{bringmann_fermi_2012,weniger_tentative_2012,su_strong_2012}. Although such a DM annihilation line should be loop-suppressed by a factor $\\sim 10^2 - 10^4$ compared to the expected continuum gamma-ray production in typical DM models, the fact that it is difficult to produce such a high energy line with astrophysical processes \\citep[however, see][]{aharonian_cold_2012} has motivated further exploration of the DM annihilation explanation. In fact, DM particle physics models do exist in which the continuum is suppressed with respect to the line emission \\citep[e.g.][]{cline_130_2012,buckley_implications_2012,bergstroem_130_2012,dudas_extra_2012,chalons_neutralino_2012}. In the analysis of \\citet{su_strong_2012}, the line's significance is maximized at Galactic coordinates $(\\ell,b) = (1.5^\\circ,0^\\circ)$, i.e. displaced from Sgr A*, the presumed dynamical center of the Galaxy, by about 200 projected parsec in the disk plane \\citep[see also][]{tempel_fermi_2012}. This offset has been viewed as a strike against a DM annihilation interpretation of the line signal. It is possible that the offset is simply due to small number statistics \\citep{yang_statistical_2012}, but nevertheless it is commonly viewed as a strike against a DM annihilation interpretation of the line signal.  On the other hand, the GC is a dynamically and energetically complex region \\citep{genzel_galactic_2010}. It harbors a supermassive black hole (SMBH) \\citep{ghez_measuring_2008,gillessen_monitoring_2009}, which may have been active as recently as 10 Myr ago, if the giant Fermi bubbles \\citep{dobler_fermi_2010,su_giant_2010} are interpreted as resulting from a black hole accretion event; it is rich in massive young stars \\citep{bartko_evidence_2009}, which likely formed in a single star burst event $6 \\pm 2$ Myr ago \\citep{paumard_two_2006}; it hosts plenty of highly energetic compact objects radiating in X-rays and gamma-rays \\citep{muno_deep_2003,muno_overabundance_2005,abazajian_consistency_2011}; and on $\\sim$kpc scales, its gravitational potential is non-axisymmetric due to the presence of a stellar bar and a boxy bulge \\citep{blitz_direct_1991,martinez-valpuesta_unifying_2011}. Given that our Galaxy is baryon-dominated inwards of $\\sim 5 - 10$ kpc \\citep{klypin_cdm-based_2002}, one might expect astrophysical processes to modify the underlying DM distribution in significant ways. In principle these processes may even displace the maximum of the DM density away from the dynamical center, thus greatly affecting the expected DM annihilation signal from the GC.  In isolated galaxy simulations, resonant interactions between the stellar bar and the DM halo have been shown to alter the shape of DM halos reducing their triaxiality \\citep{berentzen_stellar_2006,machado_loss_2010}, and to flatten a central cusp into a core \\citep{weinberg_bar-driven_2002,athanassoula_what_2003,holley-bockelmann_bar-induced_2005,weinberg_bar-halo_2007}. The latter results remain controversial, however, since other numerical studies do not see such strong effects \\citep{sellwood_bars_2003,valenzuela_secular_2003,colin_bars_2006}. Regarding an off-center DM density peak, these interactions are interesting because they may also induce a ``dark bar'' and other non-axisymmetric perturbations in the DM \\citep{athanassoula_bar-halo_2002,ceverino_resonances_2007,mcmillan_halo_2005}.  External gravitational perturbations, for example during a merger or a near passage of a satellite galaxy, could displace the tightly-bound baryonic component from the center of the overall mass distribution (the DM halo). Such offsets have been measured in galaxy clusters \\citep{allen_resolving_1998}, in which separations between the center of the X-ray emitting gas and the gravitational center determined from strong lensing can be as large as $\\sim 30$ kpc for relaxed clusters \\citep{shan_offset_2010}. A study of THINGS galaxies by \\citet{trachternach_dynamical_2008} found that offsets between the photometric and dynamical centers were less than one radio beam width ($\\sim 10''$ or $150 - 700$ pc) for 13 out of 15 galaxies with well-constrained photometric centers. However, two galaxies in their sample, NGC 3627 (a barred Sb galaxy showing signs of a recent interaction) and NGC 6946 (a barred Scd galaxy), exhibit moderate offsets between one and two beam widths.  An offset DM density peak may be a reflection of an intrinsic lopsidedness in the Galaxy \\citep{saha_milky_2009}. In fact, disk galaxies commonly exhibit substantial asymmetry in the central regions of their light distribution \\citep[for a review, see][]{jog_lopsided_2009}. The origin of these asymmetries is not fully understood, with tidal encounters, gas accretion, and a global gravitational instability being some of the physical mechanisms under consideration. The central regions of advanced galaxy mergers often show long-lived unsettled sloshing behavior \\citep{schweizer_colliding_1996,jog_measurement_2006}, and even in isolated galaxies the dynamical center can remain unrelaxed for many dynamical times \\citep{miller_off-center_1992}, especially in systems with a cored mass distribution.  Supernova- or AGN-driven gas outflows may rapidly and non-adiabatically alter the potential in the central regions of proto-galaxies, prior to the formation of the bulk of their stars. Repeated episodes of such impulsive outflows, followed by slow adiabatic re-accretion of gas, may irreversibly transfer energy to the DM, flattening the central cusp in the process \\citep{read_mass_2005,pontzen_how_2012}. Although this effect may not by itself produce an off-center DM density peak, the resulting cored density profile will be more susceptible to perturbations.  Lastly, the presence of a SMBH has been argued to lead to the formation of a steep cusp of DM \\citep{gondolo_dark_1999,gnedin_dark_2004} in the inner parsec centered on the SMBH. This would obviously preclude any significant offset between the peak in DM annihilation signal and Sgr A*.  In this work we report on the search for the presence of an offset between the dynamical center and the maximum DM density in the Eris simulation \\citep{guedes_forming_2011}, one of the highest resolution and most realistic cosmological simulations of the formation of a Milky-Way-like barred spiral galaxy. We find evidence for such an offset, at a scale of $300 - 400$ pc, consistent with the offset seen by \\citet{su_strong_2012}. Since such an offset is only seen in the dissipational hydrodynamic simulations, and not in our collisionless pure-DM runs, we suggest that baryonic physics is in some way responsible for its formation. We examine a number of different physical mechanisms, but the limited resolution of this study does not allow us to conclusively settle on a single preferred explanation. At this stage, we wish to draw attention to the possibility of the maximum DM density (and hence annihilation luminosity) not being coincident with the dynamical center of our Galaxy, commonly associated with Sgr A*. We hope that our results will stimulate future work, examining other high resolution hydrodynamic galaxy formation simulations, and investigating in more detail the physical mechanisms that can give rise to an offset DM density peak.  The remainder of this paper is organized as follows. In Section 2 we describe the numerical simulations that we have analyzed. In Section 3 we present the evidence for a DM offset in the Eris simulation. In Section 4 we go over several possible formation mechanisms and confront each of them with data from the simulations. In Section 5 we discuss implications for indirect detection searches towards the GC, and finally in Section 6 we present our conclusions. ",
        "conclusions": "We have analyzed the distribution of DM in the central regions of the hydrodynamical galaxy formation simulation Eris, one of the highest resolution and most realistic simulations to date of the formation of a barred spiral galaxy like our own Milky Way. Surprisingly, we find that the peak of the DM density in Eris is typically offset from its dynamical center by several hundred parsec. No such offset is observed in its DM-only twin simulation ErisDark, nor in the much higher resolution DM-only Via Lactea II and GHalo simulations.  The DM offset in Eris begins to appear around $z=1.5$ and grows over a period of 2 Gyr to a stable value of $\\langle \\Doff \\rangle = 340$ pc (almost three gravitational softening lengths), with a dispersion of 50 pc. The onset and duration of the DM offset appears to be well correlated with the formation of a nearly constant DM density core. The distributions of $\\rho_{\\rm max}$ and $\\Doff$ over the past 4 Gyr are inconsistent with a statistical fluctuation. Neither is the density peak a gravitationally bound structure, which rules out an incompletely disrupted subhalo core as an explanation. The most likely explanation may be a density-wave-like excitation by the stellar bar, possibly related to the resonant mechanism proposed by \\citet{weinberg_bar-driven_2002, weinberg_bar-halo_2007} to explain the transformation of a central DM cusp into a core. Arguments in favor of this explanation are the fact that the DM offset appears preferentially near the disk plane, that it is aligned to $\\sim 30$ degrees with the orientation of the stellar bar, and that is shows a periodicity of $\\sim 70$ Myr.  A central DM offset is of particularly interest in the context of the recent report by \\citet{su_strong_2012} of a highly significant detection of gamma-ray line emission from a region $\\sim 1.5$ degrees ($\\sim 200$ pc projected) away from the Galactic Center. At first impression such a large angular offset would seem to argue against a DM annihilation interpretation of this signal. Our work demonstrates that in fact just such an offset should perhaps be expected. We note, however, that at the current resolution of our numerical simulations the low contrast ($5 - 15\\%$) of the annihilation surface brightness between the offset peak and the GC may be too small to accommodate a DM annihilation explanation.  We conclude by acknowledging that properly resolving the effects of baryonic physics on the central DM distribution, in particular those involving resonant bar-halo interactions, requires much higher resolution than we have been able to afford so far in cosmological hydrodynamics simulations. Further studies at higher resolution and exploring different baryonic physics implementation are sorely needed. Of particular importance are clarifying the role of the star formation threshold parameter and supernovae feedback for the formation of a DM core, as well as the influence of the supermassive black hole at the GC on the DM distribution."
    },
    "1208/1208.1185_arXiv.txt": {
        "abstract": "We construct models of universe with a generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n})c^2$ having a linear component and a polytropic component. The linear equation of state $p=\\alpha\\rho c^2$ with $-1\\le \\alpha\\le 1$ describes radiation ($\\alpha=1/3$), pressureless matter ($\\alpha=0$), stiff matter ($\\alpha=1$), and vacuum energy ($\\alpha=-1$). The polytropic equation of state $p=k\\rho^{1+1/n} c^2$ may be due to Bose-Einstein condensates with repulsive ($k>0$) or attractive ($k<0$) self-interaction, or have another origin. In this paper, we consider the case where the density increases as the universe expands. This corresponds to a ``phantom universe'' for which $w=p/\\rho c^2<-1$ (this requires $k<0$). We complete previous investigations on this problem and analyze in detail the different possibilities. We describe the singularities using the classification of [S. Nojiri, S.D. Odintsov, S. Tsujikawa, Phys. Rev. D {\\bf 71}, 063004 (2005)]. We show that for $\\alpha>-1$ there is no Big Rip singularity although $w\\le -1$. For $n=-1$, we provide an analytical model of phantom bouncing universe ``disappearing'' at $t=0$.  We also determine the potential of the phantom scalar field and phantom tachyon field corresponding to the generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n})c^2$. ",
        "introduction": "In previous papers of this series, we have constructed models of universe with a generalized equation of state \\begin{equation} \\label{intro1} p=(\\alpha \\rho+k\\rho^{1+1/n}) c^2, \\end{equation} having a linear component and a polytropic component.  In Papers I and II, we have assumed $\\alpha+1+k\\rho^{1/n}\\ge 0$ corresponding to $w=p/\\rho c^2\\ge -1$. In that case, the density decreases as the universe expands. For $n>0$, the polytropic component dominates the linear component when the density is high: This describes the early universe (Paper I). For $n<0$, the polytropic component dominates the linear component when the density is low: This describes the late universe (Paper II).  When the polytropic pressure is positive ($k>0$), the solutions of the Friedmann equations exhibit past or future singularities (or peculiarities).  When the polytropic pressure is negative ($k<0$), there is no singularity. Furthermore, the polytropic equation of state implies the existence of an upper bound $\\rho_{max}$ (in the past) and a lower bound $\\rho_{min}$ (in the future) for the density. It makes sense to identify the maximum density to the Planck density $\\rho_P=5.16 \\, 10^{99}\\, {\\rm g}/{\\rm m}^3$ and the minimum density to the cosmological density $\\rho_{\\Lambda}=7.02\\, 10^{-24}\\, {\\rm g}/{\\rm m}^3$. These constant densities imply in turn the existence of two phases of exponential inflation, one in the early universe and one in the late universe. During the inflation, the universe is accelerating. The early inflation is necessary to solve notorious difficulties such as the singularity problem, the flatness problem, and the horizon problem \\cite{guth,linde}. The late inflation is necessary to account for the observed accelerating expansion of our universe \\cite{novae} driven by dark energy \\cite{cst}. In that context, the equation of state (\\ref{intro1}) with $k<0$ and $n<0$ corresponds to the generalized Chaplygin gas \\cite{chaplygin} that has been proposed as a model for dark energy. From the generalized polytropic equation of state (\\ref{intro1}), we have obtained a model of universe without singularity that possesses striking ``symmetries'' between the past and the future (aioniotic universe). This model, which could have been obtained from a principle of ``simplicity'' without making any observation, turns out to be strikingly consistent with what we know of the real universe. It is consistent with the standard model \\cite{weinberg,bt} but refines it by removing the primordial singularity (Big Bang). In this model, the Planck density and the cosmological density are interpreted as two {\\it fundamental} bounds for the density determined by the Planck constant $\\hbar$ (microphysics) and the cosmological constant $\\Lambda$ (cosmophysics), respectively. These bounds differ by $122$ orders of magnitudes, a difference that appears to be quite natural instead of representing a ``problem'' \\cite{weinbergcosmo}.   In this paper, we consider a case that has not been treated in our previous papers. This is the case where the density increases as the universe expands. Since the nature of dark energy is unknown, this situation cannot be rejected {\\it a priori}. It corresponds to an equation of state parameter $w$ less than $-1$ which violates the null dominant energy condition. This is referred to as a ``phantom universe'' \\cite{caldwell} because  when the equation of state with $w<-1$ is constructed in terms of a scalar field, the corresponding kinetic term has the wrong sign (negative kinetic energy). It represents therefore a phantom (ghost) scalar field (see reviews \\cite{cst,saridakis}).   Actually, there is a rich recent literature on this situation (more than one thousand papers are related to phantom dark energy) since observations do not exclude the possibility that we live in a phantom universe. Indeed, observational data indicate that the equation of state parameter $w$ lies in a narrow strip around $w=-1$ possibly being below this value \\cite{observations}.  The models based on phantom dark energy usually predict a future singularity in which the scale factor, the energy density, and the pressure of the universe become infinite in a finite time. This would lead to the death of the universe in a singularity called ``Big Smash'' \\cite{innes}, ``Big Rip'' or ``Cosmic Doomsday'' \\cite{caldwellprl}. Contrary to the ``Big Crunch'', the universe is destroyed not by excessive contraction but rather by excessive expansion.  In phantom cosmology, every gravitationally bound system ({\\it e.g.} the solar system, the Milky Way, the local group, galaxy clusters) is dissociated before the singularity \\cite{caldwellprl,np}, and the black holes gradually lose their mass and finally vanish \\cite{pedro,thermophantom}. This scenario allows the explicit calculation of the rest of the lifetime of our universe. Actually, as we approach the singularity, the energy scale may grow up to the Planck one, giving rise to a second quantum gravity era. Eventually, quantum effects may moderate or even prevent the singularity \\cite{quantum}. Other aspects of phantom cosmology have been studied in \\cite{ghosts}.  There are many interesting recent works on the study of singularities. In particular,  Nojiri {\\it et al.}  \\cite{classification} considered an equation of state of the form $p=-\\rho-f(\\rho)$ and obtained a classification of finite-time future singularities (see complements in \\cite{fernandez}). They are of four types:  $\\bullet$ Type 0 (Big Bang or Big Crunch): For $t\\rightarrow t_s$, $a\\rightarrow 0$, $\\rho\\rightarrow +\\infty$, and $|p|\\rightarrow +\\infty$.  $\\bullet$ Type I (Big Rip): For $t\\rightarrow t_s$, $a\\rightarrow +\\infty$, $\\rho\\rightarrow +\\infty$, and $|p|\\rightarrow +\\infty$.  $\\bullet$ Type II (sudden singularity): For $t\\rightarrow t_s$, $a\\rightarrow a_s$, $\\rho\\rightarrow \\rho_s$, and $|p|\\rightarrow +\\infty$.  $\\bullet$ Type III (Big Freeze): For $t\\rightarrow t_s$, $a\\rightarrow a_s$, $\\rho\\rightarrow +\\infty$, and $|p|\\rightarrow +\\infty$.  $\\bullet$ Type IV (generalized sudden singularity): For $t\\rightarrow t_s$, $a\\rightarrow a_s$, $\\rho\\rightarrow \\rho_s$, $|p|\\rightarrow p_s$, and higher derivatives of $H$ diverge\\footnote{We shall not consider this type of singularities in this paper.}.  In this classification,  $t_s$, $a_s$, $\\rho_s$, and $p_s$ are all finite constants ($a_s\\neq 0$). Type 0 is the standard Big Bang or Big Crunch singularity arising in the original Friedmann models \\cite{weinberg}. Type I is the Big Rip singularity which emerges from the phantom equation of state $p=\\alpha\\rho c^2$ with constant $\\alpha<-1$ \\cite{caldwell,caldwellprl}, and from the equation of state (\\ref{intro1}) with $\\alpha=-1$, $k<0$ and $n<-2$ \\cite{stefancic}. Type II corresponds to the sudden future singularity found by Barrow \\cite{barrow} at which $a$ and $\\rho$ are finite but $p$ diverges. Type III, arising in the equation of state (\\ref{intro1}) with $n>0$ \\cite{stefancic,bigfreeze} differs from the sudden future singularity in the sense that $\\rho$ diverges. Type IV appears in the model described in \\cite{classification}.  It is important to stress that the phantom models with $w<-1$ do not necessarily lead to future singularities. For example, the equation of state  (\\ref{intro1}) with $\\alpha=-1$ and $-2\\le n<0$ does not present future singularity \\cite{stefancic}. However, since the scale factor and the density increase indefinitely, this has been called ``Little Rip'' \\cite{littlerip}.   On the other hand, the models with $w>-1$ may lead to past or future singularities. For example, the new form of primordial singularity (for $n>0$ and $k>0$)  described in Secs. IV C, VI and in Appendix A   of Paper I corresponds to a past singularity of type III: The universe starts at $t=0$ with a finite scale factor and an infinite density. On the other hand, the future singularity (for $-1<n<0$ and $k>0$) described in Sec. IV C and in Appendix B  of Paper II corresponds to a singularity of type II: At a finite time $t_s$, the universe reaches a point at which the scale factor is finite, the density vanishes and the pressure is infinite.   In paper II, we have also introduced a notion of ``peculiarity''. This is when the density vanishes $\\rho=0$ while the scale factor has a finite value $a_s$ (when $a_s=0$ we shall call it generalized peculiarity). In that case, the universe is empty (in other works, it ``disappears''). Although there is no singularity, this situation is very peculiar. However, since the nature of dark energy is unknown, all possibilities should be contemplated.   In this paper, we provide an exhaustive study of the equation of state (\\ref{intro1}) in the case $w\\le -1$ (requiring $k<0$) for arbitrary $-1\\le \\alpha\\le 1$ and $n$. This is a natural continuation of our previous works which assumed $w\\ge -1$ (Papers I and II). Our paper also completes previous studies of the case $w\\ge -1$ that considered $\\alpha=-1$ \\cite{stefancic} or  $\\alpha=0$ \\cite{bigfreeze}. An interesting result of our study is that the equation of state (\\ref{intro1}) with $\\alpha>-1$ does {\\it not} present a Big Rip singularity although $w<-1$, contrary to the linear equation of state $p=\\alpha\\rho c^2$ with $\\alpha<-1$ \\cite{caldwellprl} or the equation of state (\\ref{intro1}) with $\\alpha=-1$ and $n<-2$ \\cite{stefancic}. Another interesting result of our study is the construction of a bouncing phantom universe for $-2<n<0$. For  $-1<n<0$, this bouncing universe presents a past singularity of type II: At $t=0$, the the pressure is infinite while the scale factor has a finite value and the density vanishes. For $\\alpha>-1$ and $n=-1$, corresponding to a constant negative pressure, the bouncing phantom universe admits a simple analytical expression.    The paper is organized as follows. In Sec. \\ref{sec_basic}, we recall the basic equations of cosmology. In Secs. \\ref{sec_ges} and \\ref{sec_dark}, we study the generalized equation of state (\\ref{intro1}) for any value of the parameters $-1<\\alpha\\le 1$, $k<0$ and $n$, assuming $w<-1$ (phantom cosmology). In Sec. \\ref{sec_scalar}, we determine the potential of the phantom scalar field  and the potential of the  phantom tachyon field corresponding to the generalized equation of state (\\ref{intro1}). In Appendix \\ref{sec_eosgm}, we treat the case $\\alpha=-1$. In Appendix \\ref{sec_summary}, we summarize all the results obtained in our series of papers and analyze the different singularities in terms of the classification of \\cite{classification}. ",
        "conclusions": "In this paper, we have performed an exhaustive study of the generalized equation of state (\\ref{basic3}) in the case where the pressure increases with the scale factor. This corresponds to the so-called phantom cosmology \\cite{caldwell}.  The case $\\alpha=-1$ was previously treated in \\cite{stefancic}. For $n<-2$, the universe experiences a future singularity of type I (Big Rip): The scale factor and the density diverge at a finite time. For $-2\\le n<0$, the scale factor and the density diverge in infinite time (Little Rip). For $n>0$, the universe experiences a future singularity of type III: The density diverges at a finite time while the scale factor tends to a constant.  We have found that when $\\alpha>-1$, the universe does not experience a Big Rip singularity. For $n<0$, there is a phase of late inflation and, for $n>0$, the universe experiences a future singularity of type III. The past evolution of the universe is interesting. For $n>0$, there is a phase of early inflation. For $n<0$, the universe exhibits a past peculiarity since the density vanishes while the scale factor tends to a finite value. For $n\\le -2$, the evolution of the scale factor is similar to the Eddington-Lema\\^itre model (the universe is static in the infinite past and grows exponentially rapidly in the future) but the evolution of the density is very different (it starts from zero in the infinite past and increases as the universe expands). For $-2<n<0$, we have obtained a model of bouncing universe which also possesses peculiar features (the density decreases in the past, vanishes at $t=0$, and increases in the future).  For $-1<n<0$, this bouncing universe presents a past singularity of type II since the pressure at $t=0$ is infinite while the scale factor is finite and the density vanishes. For $n=-1$, corresponding to a constant negative pressure, the bouncing phantom universe admits a simple analytical expression.   Of course, most of these models are academic, and do not correspond to the true evolution of our universe. However, we believe that it is important to study the equation of state (\\ref{basic3}) in full generality. On the other hand, there are indications \\cite{observations} that the equation of state parameter $w$ of our universe may become less than $-1$ in the close future (or even at present), so the {\\it late} evolution of the phantom models described in this paper may be physically relevant.   A drawback of the simple form of phantom cosmology considered in this paper is that it does not connect smoothly to the matter era (which has $w=0$). Therefore, we cannot realistically extend the phantom models to the past and obtain unified models of dust matter and phantom dark energy ($w<-1$), contrary to the unified models of dust matter and quintessence dark energy ($w>-1$) based on the generalized Chaplygin gas considered in Paper II.  A unification of dust matter and phantom dark energy can be achieved in more general models allowing to cross the phantom divide \\cite{divide}. This generalization assumes an interaction between dark matter and dark energy. These models are very interesting because they may provide a solution to the ``cosmic coincidence problem'' (the fact that the ratio of dark matter and dark energy is of order one).  The phantom cosmology is also interesting for its connection to Hoyle's version of the steady state theory \\cite{hoyle}, for its connection to wormholes \\cite{wormholes}, and for its very strange thermodynamics allowing for the existence of negative temperatures \\cite{thermophantom} like in 2D turbulence \\cite{onsager}.   However, we may recall that there is no firm evidence that we live in a phantom universe. The model of Paper II, corresponding to the standard $\\Lambda$CDM model with the primordial singularity removed, may correctly describe the whole evolution of our universe. Therefore, a more precise determination of the equation of state parameter $w$ will help discriminate between these different models.          \\appendix"
    },
    "1208/1208.1668.txt": {
        "abstract": "In this work we investigate the detectability of the gravitational stochastic background produced by cosmological sources in scenarios of structure formation. The calculation is performed in the framework of hierarchical structure formation using a Press-Schechter-like formalism. The model considers the coalescences of three kind of binary systems, namely double neutron stars (NS-NS), the neutron star-black hole binaries (NS-BH), and the black hole-black hole systems (BH-BH). We also included in the model the core-collapse supernovae leaving black holes as compact remnants. In particular, we use two different dark-energy scenarios, specifically cosmological constant ($\\Lambda$) and Chaplygin gas, in order to verify their influence on the cosmic star formation rate, the coalescence rates, and on the  gravitational wave backgrounds. We calculate the gravitational wave signals separately for each kind of source as well as we determine their collective contribution for the stochastic background of gravitational waves. Concerning to the compact binary systems, we verify that these sources produce stochastic backgrounds with signal-to-noise ratios (S/N) $\\sim 1.5$ ($\\sim 0.90$) for NS-NS, $\\sim 0.50$ ($\\sim 0.30$) for NS-BH, $\\sim 0.20$ ($\\sim 0.10$) for BH-BH, and $\\sim 0.14$ ($\\sim 0.07$) for core-collapse supernovae for a pair of advanced LIGO detectors in the cosmological constant (Chaplygin gas) cosmology. Particularly, the sensitivity of the future third generation of detectors as, for example, the Einstein Telescope (ET), in the triangular configuration, could increase the present signal-to-noise ratios by a high-factor ($\\sim 300 - 1000$) when compared to the (S/N) calculated for advanced LIGO detectors. As an example, the collective contribution of these sources can produce $({\\rm S/N})\\sim 3.3$ ($\\sim 1.8$) for the $\\Lambda$ (Chaplygin gas) cosmology for a pair of advanced LIGO interferometers and within the frequency range $\\sim 10\\,{\\rm Hz} - 1.5\\, {\\rm kHz}$. Considering ET we have $({\\rm S/N})\\sim 2200$ ($\\sim 1300$) for the $\\Lambda$ (Chaplygin gas) cosmology. Thus, the third generation of gravitational wave detectors could be used to reconstruct the history of star formation in the Universe as well as for contributing with the characterization of the dark energy, for example, identifying if there is evidence for the evolution of the dark energy equation-of-state parameter $w(a)$. ",
        "introduction": "In the previous Sections, we have analyzed the main characteristics of the stochastic backgrounds produced by compact binary systems, core-collapse to form black holes, and the composite signal of these cosmological objects. We verify that higher signal-to-noise ratios can be produced if we consider Einstein Telescope in triangular configuration. Although these cosmological sources are connected by the CSFR, we know that there are uncertainties in the minimum coalescence time-scales of NS-NS, NS-BH, and BH-BH binaries. On the other hand, the local coalescence rates of these systems can vary up to three orders of magnitude. In addition, the minimum mass able to form a black hole may vary from $\\sim 25{\\rm M}_{\\odot}$ to $\\sim 40{\\rm M}_{\\odot}$. Thus, in this Section, we present an analysis of these uncertainties and their influence on the stochastic backgrounds discussed here. We also discuss if there is a clear difference between the $\\Lambda$CDM and the Chaplygin gas which would permit constrain both the constant feature (or not) of the dark-energy equation of state and the CSFR derived for $\\Lambda$CDM and Chaplygin gas. In particular, we have analyzed:  \\begin{figure} \\includegraphics[width=90mm]{fig15.pdf} \\caption{Collective spectra of the three compact binaries taking into account the uncertainties in the parameters. The black area describes all the possible GW signals for the $\\Lambda$CDM case. The gray area represents the GW backgrounds for Chaplygin gas with $\\alpha=0.2$.} \\end{figure}  a) The local coalescence rate: This parameter acts like an offset and it does not modify the shapes of the spectra. Note, however, that ${\\rm (S/N)}\\propto \\dot\\rho_{\\rm c}^{0}(0)$ and so our values for the signal-to-noise ratios can vary from 0.1 to 10 of those listed in Tables $2-4$. Thus, this parameter can only modify the values of the (S/N) as ${\\rm (S/N)} = \\dot\\rho_{\\rm c}^{0}(0)/\\dot\\rho_{\\rm c}^{0}(0)_{u}\\times {\\rm (S/N)}_{u}$ where the subscript $u$ means the values used and derived in this work. Note that same in the worst case ($\\dot\\rho_{\\rm c}^{0}(0) =0.1\\times \\dot\\rho_{\\rm c}^{0}(0)_{u})$, it would be possible to have ${\\rm (S/N)}> 10$ (ET) for the composite signals of these three binary sources. Concerning to the duty cycle, observe that $D(z)$ is also proportional to $\\dot\\rho_{\\rm c}^{0}(0)$. Thus, the redshifts of transition from popcorn to continuous regimes (and from shot noise to popcorn regimes) can change according to the values of $\\dot\\rho_{\\rm c}^{0}(0)$;  b) The minimum stellar mass to form a black hole: This parameter basically changes the maximum frequency of the background formed by core-collapse. In the case $m_{\\rm min} = 25{\\rm M}_{\\odot}$, we obtain $\\nu_{\\rm max}=4.8\\,{\\rm kHz}$ while for $m_{\\rm min} = 40{\\rm M}_{\\odot}$ we have $\\nu_{\\rm max}= 1.2\\,{\\rm kHz}$. In terms of (S/N), if we change $m_{\\rm min}$ from $40{\\rm M}_{\\odot}$ to $25{\\rm M}_{\\odot}$ the signal-to-noise ratios increase in $5\\%$ in relation to those values present in Table 5. Looking for the results in Table 7, collective contribution of all sources, (S/N) increases by $2\\%$ ($50\\%$) for Adv. LIGO (ET). The frequency where $\\Omega_{\\rm GW}$ peaks is weakly dependent on this parameter in both cases $\\Lambda$CDM and Chaplygin gas;  c) Efficiency of generation of gravitational waves ($\\varepsilon_{\\rm GW}$): Note that $\\Omega_{\\rm GW} \\propto \\varepsilon_{\\rm GW}$. Thus, same with an efficiency of generation of GWs $\\sim 10^{-5}-10^{-6}$ could be possible to have $({\\rm S/N})> 10$ for ET in triangular configuration (see Table 5);  \\begin{figure} \\includegraphics[width=90mm]{fig16.pdf} \\caption{Collective spectra of all sources studied in this work and taking into account the uncertainties in the parameters. The black area describes all the possible GW signals for the $\\Lambda$CDM case. The gray area represents the GW backgrounds for Chaplygin gas with $\\alpha=0.2$.} \\end{figure}  d) Coalescence time-scale of NS-NS: We change this parameter from $20\\,{\\rm Myr}$ to $100\\,{\\rm Myr}$. As a consequence, the coalescence rate peaks at $z\\sim 1.9\\, (1.30)$ instead of $z\\sim 2.27\\, (1.63)$ for the $\\Lambda$CDM (Chaplygin gas with $\\alpha = 0.2$) while the (S/N) of the Table 2 typically decreases by $15\\%$. Looking for the collective contribution of all sources in Table 7 we note that (S/N) decreases by $8\\%$ ($3\\%$) for Adv. LIGO (ET). There is just a slight modification of the frequency where $\\Omega_{\\rm GW}$ peaks.  e) Coalescence time-scale of NS-BH: We change this parameter from $10\\,{\\rm Myr}$ to $50\\,{\\rm Myr}$. As a consequence, the coalescence rate peaks at $z\\sim 2.1\\, (1.50)$ instead of $z\\sim 2.45\\, (1.81)$ for the $\\Lambda$CDM (Chaplygin gas with $\\alpha = 0.2$) while the (S/N) of the Table 3 typically decreases by $10\\%$. Looking for the collective contribution of all sources in Table 7 we note that (S/N) decreases by $1\\%$ for both Adv. LIGO and ET.  f) Coalescence time-scale of BH-BH: We change this parameter from $100\\,{\\rm Myr}$ to $500\\,{\\rm Myr}$. As a consequence, the coalescence rate peaks at $z\\sim 1.25\\, (0.83)$ instead of $z\\sim 1.86\\, (1.31)$ for the $\\Lambda$CDM (Chaplygin gas with $\\alpha = 0.2$) while the (S/N) of the Table 4 typically decreases by $20\\%$. Looking for the collective contribution of all sources in Table 7 we note that (S/N) decreases by $0.8\\%$ for both Adv. LIGO and ET.  A question could arise about the uncertainties described above: Is it possible to have a clear separation of the two backgrounds ($\\Lambda$CDM and Chaplygin gas cosmologies) or the uncertainties listed above produce a superposition of these backgrounds? A second question could also arise: Can different dark energy scenarios produce distinct signatures on the CSFR? In order to answer these questions, we present in Figures 15 and 16 the gravitational wave backgrounds with the uncertainties discussed above and for the models A1 and A3 of Table 1. In these Figures, we just keep fixed two parameters: $\\varepsilon_{\\rm GW}=10^{-4}$ and $m_{\\rm min}=40\\,{\\rm M}_{\\odot}$. Below $\\nu_{\\rm obs}\\sim 1\\, {\\rm kHz}$ there is no superposition between the GW signals in the case $\\Lambda$CDM and Chaplygin gas ($\\alpha = 0.2$), same with all the uncertainties in the parameters. However, note that the case Chaplygin gas with $\\alpha =1$ can not be separated from $\\alpha=0.2$. There is a superposition between these two Chaplygin models if we take into account all the uncertainties discussed above.  \\begin{figure} \\includegraphics[width=90mm]{fig17.pdf} \\caption{Possible CSFRs taking into account all the viable models studied in this work. The black are represents the family of CSFRs for the $\\Lambda$CDM cosmology (models A1 to A3) while the gray area shows the family of CSFRs for the Chaplygin gas (models A4 to A6) as dark-energy component of the Universe. Note that there is no overlap at $z> 2$ between these two dark-fluids same with all the uncertainties in the parameters.} \\end{figure}  The second point is related to the CSFR. Looking for Figure 17, we see that there is no overlap between the case $\\Lambda$CDM and Chaplygin gas at $z>2$ same with all the uncertainties in the parameters. The areas defined by $\\Lambda$CDM and Chaplygin gas cosmology do not overlap in the redshift range $[2-20]$. In principle, with observational data less scattered in the range $z\\sim 2-5$, it would be possible to have a better indication of the dark-energy equation of state from the observed CSFR. Otherwise, being detected a stochastic background of GWs with high (S/N), as in the case of ET, we could work with the inverse problem reconstructing the CSFR from the observed background. In this way, ET could contribute with a better comprehension of how star formation is regulated at high redshifts. ",
        "conclusions": ""
    },
    "1208/1208.5904_arXiv.txt": {
        "abstract": "We study a solar flare that occurred on September 10, 2002, in active region NOAA 10105 starting around 14:52 UT and lasting approximately 5 minutes in the radio range. The event was classified as M2.9 in X-rays and 1N in H$\\alpha$. Solar Submillimeter Telescope observations, in addition to microwave data give us a good spectral coverage between 1.415 and 212 GHz. We combine these data with ultraviolet images, hard and soft X-rays observations and full-disk magnetograms. Images obtained from Ramaty High Energy Solar Spectroscopic Imaging data are used to identify the locations of X-ray sources at different energies and to determine the X-ray spectrum, while ultra violet images allow us to characterize the coronal flaring region. The magnetic field evolution of the active region is analyzed using Michelson Doppler Imager magnetograms. The burst is detected at all available radio-frequencies. X-ray images (between 12~keV and 300~keV) reveal two compact sources and 212 GHz data, used to estimate the radio source position, show a single compact source displaced by 25\\arcsec\\ from one of the hard X-ray footpoints. We model the radio spectra using two homogeneous sources, and combine this analysis with that of hard X-rays to understand the dynamics of the particles.  Relativistic particles, observed at radio wavelengths above 50~GHz, have an electron index evolving with the typical {\\em soft--hard--soft} behaviour. ",
        "introduction": "High frequency radio observations, above 50~GHz,  bring information about relativistic particles (see e.g. \\opencite{Ramatyetal:1994} and \\opencite{Trottetetal:1998}). Moreover, the efficiency of synchrotron emission, responsible for the radio radiation, increases as the electron energy increases, contrary to the bremsstrahlung mechanism which is the origin of the Hard X-ray (HXR) emission \\cite{Whiteetal:2011}. This makes observations at high frequencies very attractive for the analysis of high energy particles.  For typical magnetic fields on the Chromosphere and mildly relativistic electrons, gyrosynchrotron theory expects a peak frequency at approximately 10~GHz. Therefore the caveat of submillimeter observations is that flare emission becomes weaker as the observing frequency increases. At the same time, at high frequencies, earth atmosphere becomes brighter and absorbs much of the incoming radiation. Notwithstanding some X-class flares have shown a second spectrum besides the microwaves spectrum,  with an optically thick emission at submillimeter frequencies, sometimes described as an {\\em upturn} (see e.g. \\opencite{Kaufmannetal:2004}, \\opencite{Silvaetal:2007}, \\opencite{Luthietal:2004b}). Nonetheless, \\inlinecite{Cristianietal:2008} found, in a medium size flare, a second radio component peaking around 200~GHz. We call these cases double radio spectrum events.\\\\  Although different mechanisms were proposed to explain the double radio spectrum events \\cite{KaufmannRaulin:2006,FleishmanKontar:2010}, the conservative approach of two distinct synchrotron sources can fit reasonably well to the observations \\cite{Silvaetal:2007,Trottetetal:2008}. We note, however, that observations at higher frequencies are needed to completely determine the radiation mechanism of those events that only show the optically thick emission of the second component, like in \\inlinecite{Kaufmannetal:2004} and, because of their strong fluxes, have much stringent requirements.\\\\  Therefore, the double radio spectrum bursts may represent a kind of events whose low frequency component is the classical gyrosynchrotron from mildly relativistic particles peaking around 10~GHz, and the high frequency component is also synchrotron emission with peak frequency around or above 50~GHz, depending on the flare characteristics (in some cases above 400~GHz).\\\\  In this work we present a detailed analysis of a double radio spectrum burst occurred during a GOES M class event on September 10, 2002 and observed in radio from 1.415 to 212~GHz.  We'll show that the low frequency component is well correlated with the HXR observed with RHESSI up to 300~keV, hence we can study the dynamics of the mildly relativistic electrons inside the coronal loop. On the other hand, the high frequency component, which is also well represented by an electron synchrotron source, should be produced by a different particle population and, likely, in a different place. \\\\  We first present the data analyzed, explain the reduction methods and give the clues that justify the interpretations in Section \\ref{sec:observations}. The spectral analysis, both at radio and Hard X-rays is the kernel of our work, thus it deserves the entire Section \\ref{sec:analysis}. We divide the interpretation of the event in two wide energy bands for the mildly relativistic and the relativistic particles in Section \\ref{sec:discusion}. The consequences of our interpretations are presented as our final remarks in Section \\ref{sec:fin}. ",
        "conclusions": "\\label{sec:discusion}  The optically thin emission above 50~GHz is produced by relativistic particles, while the optically thick emission at microwaves and the HXR observed by RHESSI are emitted by mildly relativistic particles.  We roughly divide the analysis in these two energy bands.  \\subsection{Dynamics of mildly relativistic electrons} \\label{sec:mild-relativistic}  In order to understand the dynamics of the mildly relativistic electrons we compare the HXR emission, which it was observed up to 300~keV with the {\\em low frequency} radio data. The comparison of the temporal evolution of both sets of data (see Figures \\ref{fig:profiles} and \\ref{fig:delta}) supports the existence of trapped electrons because: 1) the duration of the impulsive phase in HXR is shorter than in radio and 2) the peak time in HXR occurs before the radio peak, even at low frequencies. Therefore, the HXR time profile is not necessarily the representation of the injected electrons, since there are transport effects along the loop, or at least, HXR may represent the injected electrons that precipitate directly, without being subject to trapping.\\\\  We can use the spectral analysis to derive the rate of the injected electrons in the emitting area in function of time. To do so, we write a simplified continuity equation that depends only on time, since we are not interested on how the electron distribution changes in terms of energy, pitch angle, or depth. In this simplified model, we are interested only in the total instantaneous number of electrons, $N(t)$, inside the magnetic loop, incremented by a source, $Q(t)$, from the acceleration site and decremented by the precipitated electrons, $P(t)$. Therefore, the continuity equation should be \\begin{equation} \\frac{dN(t)}{dt} = Q(t) - P(t) \\ . \\label{eq:cont} \\end{equation} Integrating the above equation in the interval $(t_i,t_{i+1})$ (with $i=0,1,2\\dots$) and solving for $Q(t_i)$ yields \\begin{equation} Q(t_i)\\Delta t_{i+1} = N(t_{i+1}) - N(t_i) + P(t_i)\\Delta t_{i+1} \\ , \\label{eq:injection} \\end{equation} with $\\Delta t_{i+1} = t_{i+1} - t_i$.  Since we are comparing $< 300$~keV emission with gyrosynchrotron, we can identify the instantaneous number of electrons inside the loop $N(t)$ with the trapped particles emitting the lower frequency component.  On the other hand, the particles leaving the volume $P(t)$ produce the HXR emission observed by RHESSI. In our picture, the low frequency component is produced all along a loop with a length of $10^9$~cm, while the HXR emission is produced in a narrow slab (see {\\it e.g.} \\opencite{Holmanetal:2011}) with a very small surface (see Figure~\\ref{fig:MDI-HXR-sst}); therefore, we can neglect the gyroemission produced within this small volume. Furthermore, no change would be appreciated if one includes the particles responsible for the high frequency component since they are two orders of magnitude less than those that produce the low frequency component. (See Table \\ref{tbl:parametros}) \\\\  We divided the event in 10 second intervals and assumed that within these intervals the conditions do not change. The computation of Equation~\\ref{eq:injection} is straightforward and the result is presented in Figure \\ref{fig:injection}.  Since $> 100$~keV data have good S/N ratio only between 14:52:50~UT and 14:53:30~UT, (see Figure \\ref{fig:xspec-fit}) the analysis is restricted to this interval, although is clear from the time profiles that there are emitting particles before and after. We observe a continuous injection with two peaks, one at the beginning and the second during the decay of the impulsive phase. To verify our results, we sum the precipitated electrons, $\\sum_i P(t_i)\\Delta t=3.4\\times 10^{37}$, and we compare this number with the maximum number of electrons existing instantaneously inside the loop, $\\max[N(t_i)] = 4.7\\times 10^{37}$. The difference between these two numbers maybe due to the fact that we are limited to the time interval in which the HXR data are statistically meaningful; therefore, we cannot track the precipitation until the end of the gyrosynchrotron emission.\\\\  \\begin{figure}[h!] \\centerline{\\resizebox{!}{0.5\\textheight}{\\includegraphics{injection.eps}}} \\caption{Top: time profiles at 15.4 (blue) and 75--300~keV (red). Bottom: the precipitated electrons $P(t)\\Delta t$ (green); the instantaneous number of electrons, $N(t)$ (blue), and the injected electrons $Q(t)\\Delta t$ (red).} \\label{fig:injection} \\end{figure}   We observe that changes in $\\deltalf$ (Fig. \\ref{fig:delta}) can be related to the injections occurred at around 14:52:51~UT and 14:53:15~UT (Figure \\ref{fig:injection}).  The electron index $\\delta_X$ lies around 5, as $\\deltalf$ for the same period, although, it increases slowly until 14:53:20~UT  when it suddenly softens by around 0.5. The difference in time evolution of $\\delta_X$ and $\\deltalf$ is an indication that the softening of the former is a consequence of the trapping.  Since $\\deltalf$ has two constant values, we conclude that it is not affected by the medium. In the above analysis we rely upon the parameters derived from the radio and HXR data fittings.  Although we do not claim that the obtained solutions are unique, the fact that two independent fittings give very comparable results, give us confidence on them.  \\\\  The progressive delay of the radio emission observed in Figure \\ref{fig:delays} must be interpreted differently depending on the frequency range.  Between 1 and 5 GHz, the short pulse dominates the emission during the rising phase; therefore, its contribution should be removed to asses the delay of the main emission. This may lead to ambiguous results, hence, we preferred not to analyze delays in this range. In the range between 9 and 30~GHz lies the peak frequency, i.e. the optical opacity $\\tau$ is approximately 1, hence, the gyrosynchrotron self absorption is critical.  The peak frequency ($\\tau\\simeq 1$) of the low frequency component shifts from approximately 7~GHz to 15~GHz from 14:52:40~UT to the peak time around 14:5259~UT (Fig. \\ref{fig:rspec-fit}).  Since there is no change in the magnetic field, this can be interpreted by the accumulation of accelerated electrons inside the loop due to the trapping that increases its density, which is the dominant factor of the self-absorption mechanism. The shift makes the low frequencies more absorbed and thus increases the relative importance of higher frequencies.  Therefore, even with a rather constant $\\deltalf$ (Fig. \\ref{fig:delta}) the progressive delay should be observed. In our fittings at 14:52:40~UT we have $\\tau(\\nu=7\\ \\mbox{GHz})=0.7$ and $\\tau(\\nu=10\\ \\mbox{GHz})=0.15$. Later on, during peak time $\\tau(\\nu=7\\ \\mbox{GHz})=60$ and $\\tau(\\nu=10\\ \\mbox{GHz})=3.7$. We note that the emission is proportional to $1-e^{-\\tau}$, hence, we have a relative amplification of $(1-e^{-60})/(1-e^{-0.7}) \\simeq 1.4$ at 7~GHz, while it is $(1-e^{-3.7})/(1-e^{-0.15})\\simeq 7.5$ at 10~GHz. The relative amplification changes the rate at which the emission rises and, therefore, the time when the signal reaches a certain level with respect to its maximum.\\\\  \\subsection{Relativistic Electrons} \\label{sec:relativistic}  The centroid position of the emitting source at 212~GHz remains quite stable during the flare which may imply that the source is compact. Moreover, it is placed 25\\arcsec~far from one of the HXR footpoints (Figure \\ref{fig:MDI-HXR-sst}). A similar result was obtained by \\inlinecite{Trottetetal:2008} for the {\\tt SOL2003-10-28T11:10} flare; during the impulsive phase (interval B in their work) the centroid positions of the 210~GHz emission lie at approximately 10\\arcsec\\ from the center of one of the  HXR footpoints (250--300~keV), but are coincident with the location of precipitating high energy protons with energies above 30~MeV seen in $\\gamma$-ray imaging of the 2.2~MeV line emission. Since for our work we do not have $\\gamma$-ray imaging to compare with, we should be cautious because the uncertainty in position is of the order of the position shift. \\\\  Since we observe the optically thin part of the high frequency spectra, the obtained $\\deltahf$ is not affected by the medium and, within data uncertainties, it must be correct.  On the other hand, we assumed a standard viewing angle $\\theta=45^\\circ$ which gives us a mean value of the magnetic field $B$ and total number of electrons $N$.  Increasing (reducing) $\\theta$ results in smaller (larger) $B$ and $N$. Although we cannot rightly evaluate $\\theta$ with our data, we do not expect an extreme value for it since the AR is located not far from Sun center (E43). Furthermore we did consider an isotropic electron distribution. The total number of accelerated electrons is $8\\times 10^{34}$ during peak time, and they should not produce enough bremsstrahlung flux to be detected by RHESSI detectors.  We confirmed this by computing its HXR emission using the {\\tt bremthick}\\footnote{Developed by G. Holman, last revision May 2002. Obtained from the RHESSI site: \\url{http://hesperia.usfc.nasa.gov/hessi/modelware.htm}} program (magenta curves in Figure \\ref{fig:xspec-fit}) which is orders of magnitude smaller than the mildly relativistic electron emission and remains below (except during one time interval) the background. \\\\  The spectral index of the high frequency component $\\alpha_{hf}$ shows a SHS behaviour, but in this case  we cannot conclude whether its origin comes from the acceleration mechanism or from the interaction with the medium as before. We tend to think that the former should be the cause, since these are relativistic particles and their interaction with the medium should be less effective.  The emission must come from a compact region with a strong magnetic field and the electron index $\\deltahf$ should be harder than $\\delta_X$ and $\\deltalf$, as is the case. These arguments support the evidence of the existence of a separated source where relativistic electrons are the responsible for the emission.\\\\  The progressive delay of the radio emission above 50~GHz (Figure \\ref{fig:delays}) can be interpreted considering the initial hardening of the spectral index $\\alpha_{hf}$. If it is due to the acceleration mechanism that accelerates first the {\\em low} energy particles and later the {\\em high} energy particles, then, the progressive delay is a consequence."
    },
    "1208/1208.2009_arXiv.txt": {
        "abstract": "We measure apparent velocities ($v_{\\rm app}$) of absorption lines for 36 white dwarfs (WDs) with helium-dominated atmospheres -- 16 DBAs and 20 DBs -- using optical spectra taken for the European Southern Observatory SN Ia progenitor survey (SPY).   We find a difference of $6.9\\pm6.9$\\,km~s$^{-1}$ in the average apparent velocity of the H$\\alpha$ lines versus that of the \\ion{He}{1} 5876\\,\\AA\\ for our DBAs.  This is a measure of the blueshift of this He line due to pressure effects.  By using this as a correction, we extend the gravitational redshift method employed by \\citet{Falcon10} to use the apparent velocity of the \\ion{He}{1} 5876\\,\\AA\\ line and conduct the first gravitational redshift investigation of a group of WDs without visible hydrogen lines. We use biweight estimators to find an average apparent velocity, $\\langle v_{\\rm app}\\rangle_{\\rm BI}$, (and hence average gravitational redshift, $\\langle v_{\\rm g}\\rangle_{\\rm BI}$) for our WDs; from that we derive an average mass, $\\langle M\\rangle_{\\rm BI}$.  For the DBAs, we find $\\langle v_{\\rm app}\\rangle_{\\rm BI}=40.8\\pm4.7$\\,km~s$^{-1}$ and derive $\\langle M\\rangle_{\\rm BI}=0.71^{+0.04}_{-0.05}\\,M_\\odot$. Though different from $\\langle v_{\\rm app}\\rangle$ of DAs (32.57 km s$^{-1}$) at the 91\\% confidence level and suggestive of a larger DBA mean mass than that for normal DAs derived using the same method \\citep[$0.647^{+0.013}_{-0.014}\\,M_\\odot$;][]{Falcon10}, we do not claim this as a stringent detection.  Rather, we emphasize that the difference between $\\langle v_{\\rm app}\\rangle_{\\rm BI}$ of the DBAs and $\\langle v_{\\rm app}\\rangle$ of normal DAs is no larger than 9.2\\,km s$^{-1}$, at the 95\\% confidence level; this corresponds to roughly $0.10\\,M_\\odot$.  For the DBs, we find $\\langle v_{\\rm app}^{\\rm He}\\rangle_{\\rm BI}=42.9\\pm8.49$\\,km s$^{-1}$ after applying the blueshift correction and determine $\\langle M\\rangle_{\\rm BI}= 0.74^{+0.08}_{-0.09}\\,M_\\odot$.  The difference between $\\langle v_{\\rm app}^{\\rm He}\\rangle_{\\rm BI}$ of the DBs and $\\langle v_{\\rm app}\\rangle$ of DAs is $\\leq11.5$\\,km s$^{-1}$ ($\\sim0.12\\,M_\\odot$), at the 95\\% confidence level.  The gravitational redshift method indicates much larger mean masses than the spectroscopic determinations of the same sample by \\citet{Voss07}.  Given the small sample sizes, it is possible that systematic uncertainties are skewing our results due to the potential of kinematic substructures that may not average out.  We estimate this to be unlikely, but a larger sample size is necessary to rule out these systematics. ",
        "introduction": "In \\citet{Falcon10}, we show that the gravitational redshift method is an effective tool for measuring mean masses of groups of white dwarfs (WDs) and has the advantage of being mostly independent from the spectroscopic method \\citep[e.g.,][]{Bergeron92b}.  \\citet{Falcon10} investigate normal DAs, the largest class of WDs.  The next logical step is to ask whether the method can be applied to the second largest class, DBs, which constitute 20\\% of all WDs below $T_{\\rm eff}\\sim17,000$\\,K\\, and $\\sim9$\\% of WDs at higher temperatures \\citep{Beauchamp96,Bergeron11}.  However, the gravitational redshift method historically uses the apparent velocities of hydrogen Balmer line cores.  The work by \\citet{Shipman76} and by \\citet{Grabowski87} show H$\\alpha$ to be suitable for this purpose since it are not significantly affected by pressure shifts.  Pure DB spectra exhibit only helium lines, and as \\citet{Greenstein67} first pointed out, using WD photospheric helium lines for gravitational redshift measurements can be difficult due to the likelihood of systematics introduced by pressure effects. These systematics include that, in theory, and with some experimental support \\citep[e.g.,][]{Berg62,Perez03}, different helium lines can be pressure shifted by different amounts, in different (blue or red) directions, and with a dependency on temperature \\citep[e.g.,][]{Griem62,Bassalo76,Dimitrijevic90,Omar06}.  For this reason, attempts at gravitational redshift measurements for helium-dominated WDs have been sparse.  \\citet{Koester87} measures line shifts of \\ion{He}{1} 4026, 4471, 4713, and 4922\\,\\AA\\ in the spectrum of the common proper motion star WD~0615-591 and of \\ion{He}{1} 4471 and 4922\\,\\AA\\ in the wide binary WD~2129+000.  These line shifts are negative (blue), meaning that the magnitude of the pressure effects are larger than the magnitude of the gravitational redshift, which, in this case, are opposing each other.  The fact that these WDs are relatively cool \\citep{Bergeron11} is consistent with the expectation that pressure effects should be significant; we will elaborate on this point in Section \\ref{heline}.  \\citet{Koester87} concludes that due to the state of the theory at the time -- laboratory measurements and theoretical predictions often disagreed on the magnitude and sometimes even sign of the shift -- he cannot deduce meaningful gravitational redshifts for these two DBs.   \\citet{Wegner89} measure the gravitational redshift for the Hyades DBA WD~0437+138 using H$\\alpha$ and mention that the velocity ``...is unaffected by the pressure shift problems of helium'' while providing no further detail.  The importance of van der Waals broadening in cool helium-dominated atmospheres \\citep{Bergeron91}, however, was perhaps not yet well-established.  In hindsight, it is likely that H$\\alpha$ in this WD {\\it is} significantly affected.  The sample of common proper motion binary and cluster WDs in \\citet{Reid96} contains three targets with helium-dominated atmospheres. For both DBAs, WD~0437+138 and WD~1425+540, Reid determines different gravitational redshifts from using H$\\alpha$ than from H$\\beta$. He mentions that this discrepancy could be because of pressure shifts due to the high atmospheric helium abundance and deems the H$\\alpha$ result as the better redshift estimate since this line should be less affected than H$\\beta$.  Reid also measures, like \\citet{Koester87}, negative shifts of helium lines for the DB WD~2129+000.  One of his measured lines (\\ion{He}{1} 4921, 5015, and 6678\\,\\AA) is in common with that of \\citet{Koester87}.  With more recent high-resolution spectroscopic observations of helium-atmosphere WDs \\citep{Voss07} and with the method of \\citet{Falcon10}, we now have the tools to revisit the gravitational redshift of DBs.  Such an investigation is a valuable check to the latest spectroscopic work \\citep{Bergeron11}, the analysis of which is nontrivial due in part to the challenge of interpreting pressure-broadened helium lines \\citep{Beauchamp97,Beauchamp98,Beauchamp99}.  In Section \\ref{heline} we discuss using the apparent velocity of the \\ion{He}{1} 5876\\,\\AA\\ line in the context of our work and, as \\citet{Wegner87} suggest, check for consistency within the DBA sample of the apparent velocities of both the hydrogen and helium line species.  In Section \\ref{DBAs} we perform the original gravitational redshift method that uses hydrogen Balmer lines on the DBAs, which results in the most direct comparison with the DAs from \\citet{Falcon10}. Then in Section \\ref{DBs} we extend the method to WDs with only helium lines. ",
        "conclusions": "We measure the apparent velocity ($v_{\\rm app}$) of the \\ion{He}{1} 5876\\,\\AA\\ line for a sample of DBAs and compare it to that of the H$\\alpha$ Balmer line.  We find a difference of $6.9\\pm6.9$\\,km~s$^{-1}$ in the average apparent velocities from the two line species, which we attribute to the blueshift of this He line due to pressure effects \\citep[e.g.,][]{Dimitrijevic90} averaged over the sample.  With assumptions one can apply this measured blueshift as a correction to other samples of helium-atmosphere WDs.  We do so in order to investigate the average gravitational redshift of a sample of DBs for the first time.  Following the gravitational redshift method from \\citet{Falcon10}, but using biweight estimators which are better suited for small sample sizes, we find $\\langle v_{\\rm g}\\rangle_{\\rm BI}= \\langle v_{\\rm app}\\rangle_{\\rm BI}= 40.8\\pm4.7$\\,km~s$^{-1}$ for 16 DBAs with $T_{\\rm eff}\\ge16,500$\\,K from SPY.  We translate this $\\langle v_{\\rm app}\\rangle_{\\rm BI}$ to a mass: $\\langle M\\rangle_{\\rm BI} = 0.71^{+0.04}_{-0.05}\\,M_\\odot$.  Though different from the $\\langle v_{\\rm app}\\rangle$ of DAs, 32.57 km s$^{-1}$, at the 91\\% confidence level and suggestive of a larger DBA mean mass than DA, $0.647^{+0.013}_{-0.014}\\,M_\\odot$ \\citep{Falcon10}, this is not a stringent result.  We emphasize that the difference between $\\langle v_{\\rm app}\\rangle_{\\rm BI}$ of the DBAs and $\\langle v_{\\rm app}\\rangle$ of normal DAs is no larger than 9.2\\,km~s$^{-1}$, at the 95\\% confidence level; this corresponds to roughly $0.10\\,M_\\odot$.  Our $\\langle M\\rangle_{\\rm BI}$ for DBAs is also larger than the average of the spectroscopic mass determinations for these targets from \\citet{Voss07} at $0.62\\pm0.02\\,M_\\odot$.  It agrees with the average mass of the 22 DBAs with $T_{\\rm eff}\\ge16,500$\\,K from \\citet{Bergeron11} at $0.67\\pm0.02\\,M_\\odot$.  We use the \\ion{He}{1} 5876\\,\\AA\\ line to conduct the first gravitational redshift investigation of a group of WDs without visible hydrogen lines. For 20 DBs from SPY, we find $\\langle v_{\\rm app}^{\\rm He}\\rangle_{\\rm BI}=42.9\\pm8.9$\\,km s$^{-1}$ after applying the correction for our measured blueshift due to pressure effects.  We determine $\\langle M\\rangle_{\\rm BI}= 0.74^{+0.08}_{-0.09}\\,M_\\odot$.  The difference between $\\langle v_{\\rm app}^{\\rm He}\\rangle_{\\rm BI}$ of the DBs and $\\langle v_{\\rm app}\\rangle$ of DAs is $\\leq11.5$\\,km s$^{-1}$, at the 95\\% confidence level; this corresponds to roughly $0.12\\,M_\\odot$. The $\\langle M\\rangle_{\\rm BI}$ is much larger than the average of the spectroscopic mass determinations from SPY, $0.58\\pm0.02\\,M_\\odot$.  It is slightly larger than the average spectroscopic mass of DBs with $T_{\\rm eff}>16,500$\\,K from \\citet{Bergeron11}, $0.63\\pm0.01\\,M_\\odot$, and the mean of DBs with $T_{\\rm eff}>16,000$\\,K from SDSS, $0.646\\pm0.006\\,M_\\odot$ \\citep{Kepler10}.  Combining our DBA and DB samples to group all WDs with helium-dominated atmospheres, we find $\\langle v_{\\rm app}^{\\rm He}\\rangle_{\\rm BI}= 43.4\\pm7.9$\\,km s$^{-1}$ (after the correction) and determine $\\langle M\\rangle_{\\rm BI}= 0.74^{+0.07}_{-0.08}\\,M_\\odot$.  This differs from the $\\langle v_{\\rm app}\\rangle$ of DAs at the 82\\% confidence level.  The difference between $\\langle v_{\\rm app}^{\\rm He}\\rangle_{\\rm BI}$ of the DBAs + DBs and $\\langle v_{\\rm app}\\rangle$ of DAs is $\\leq7.7$\\,km~s$^{-1}$ (roughly $0.08\\,M_\\odot$), at the 95\\% confidence level.  Our $\\langle M\\rangle_{\\rm BI}$ is much larger than the average spectroscopic mass of these targets from SPY at $0.60\\pm0.02\\,M_\\odot$ and slightly larger than the average spectroscopic mass of the DBAs and DBs with $T_{\\rm eff}\\ge16,500$\\,K from \\citet{Bergeron11}, $0.65\\pm0.01\\,M_\\odot$.  Given the small sample sizes, it is possible that system uncertainties are skewing our results due to the potential of kinematic substructures that may not average out.  We estimate this to be unlikely, but a larger sample size is necessary to rule out these systematics."
    },
    "1208/1208.1907_arXiv.txt": {
        "abstract": "{There are a number of faint compact infrared excess sources in the central stellar cluster of the Milky Way. Their nature and origin is unclear. In addition to several isolated objects of this kind there is a small but dense cluster of comoving sources (IRS13N) located $\\sim$3'' west of SgrA* just 0.5'' north of the bright IRS13E cluster of Wolf-Rayet and O-type stars. Based on the analysis of their color and brightness, there are two main possibilities: (1) they may be dust-embedded stars older than a few Myr, or (2) very young, dusty stars with ages younger than 1~Myr.} {We present a first K$_s$-band identification and proper motions of the IRS13N members, the high-velocity dusty S-cluster object (DSO, also referred to as G2), and other infrared excess sources in the central field. Goal is to constrain the nature of these source. } {The $L'$- (3.8$\\mu$m) K$_s$- (2.2$\\mu$m) and H-band (1.65$\\mu$m) observations were  carried out using the NACO adaptive optics system at the ESO VLT. Proper motions were obtained by linear fitting of the stellar positions extracted by StarFinder as a function of time, weighted by positional uncertainties, and by Gaussian fitting from high-pass filtered and deconvolved images. We also present results of near-infrared (NIR) H- and K$_s$-band ESO-SINFONI integral field spectroscopy of the Galactic Center cluster ISR13N.} { We show that within the uncertainties, the positions and proper motions of the IRS13N sources in  K$_s$- and $L'$-band are identical. The HK$-s$L' colors then indicate that the bright $L'$-band IRS13N sources are indeed dust-enshrouded stars rather than core-less dust clouds. The proper motions also show that the IRS13N sources are not strongly gravitationally bound to each other. Combined with their NIR colors, this implies that they have been formed recently. For the DSO we obtain proper motions and a K$_s$-$L'$-color. } {Most of the compact $L'$-band excess emission sources have a compact H- or K$_s$-band counterpart and therefore are likely stars with dust shells or disks. Our new results and orbital analysis from our previous work favor the hypothesis that the infrared excess IRS13N members and other dusty sources close to SgrA* are young dusty stars and that star formation at the Galactic Center (GC) is a continuously ongoing process. For the DSO the color information indicates that it may be a dust cloud or a dust-embedded star.} ",
        "introduction": "\\label{intro}  The formation of young and massive stars is an important process at the center of the Milky Way (see e.$\\,$g. Paumard et al. 2006, Ghez et al. 2005). In the central half parsec of the Milky Way these objects are organized in at least one disk-like structure of clockwise rotating stars (CWS; Genzel et al. 2003, Levin \\& Beloborodov 2003, Paumard et al. 2006) and the existence of a second, less populated disk of counter-clockwise rotating stars (CCWS) has been proposed by Paumard et al. (2006). It is unclear how these young stars have been formed in the strong tidal field of the super-massive black hole (SMBH) at the position of SgrA*. There are two prominent scenarios: One includes star formation in-situ (in an accretion disk; Levin \\& Beloborodov 2003, Nayakshin et al. 2006), the other requires the in-spiral of a massive stellar cluster that formed at a safe distance of 5-30 pc from the Galactic Center (Gerhard et al. 2001, McMillan \\& Portegies Zwart 2003, Kim et al. 2004, Portegies Zwart et al. 2006). Currently the in-situ scenario appears to be favored by a number of authors (Nayakshin \\& Sunyaev 2005, Nayakshin et al. 2006, Paumard et al. 2006). Also, results by Stolte et al. (2007, 2010) seem to rule out the possibility that known compact clusters close to the GC (such as the Arches cluster) could migrate inwards and fuel the young stellar population at the very center.  A prime object for testing these scenarios both observationally and theoretically is the IRS13 group of sources. IRS13E (located $\\sim$$\\,$3'' west and $\\sim$$\\,$1.5'' south of SgrA*) is the densest stellar association after the immediate vicinity of SgrA* and contains several massive Wolf-Rayet (WR) and O-type stars (Maillard et al. 2004, Moultaka et al. 2005, Paumard et al. 2006, Fritz et al. 2010). It is generally considered to be associated with the mini-spiral (Moultaka et al. 2005, Paumard et al. 2004). For IRS13E four out of seven identified stars show highly correlated velocities (Maillard et al. 2004, Sch\\\"odel et al 2005),indicating that is probably bound. It is unclear what the origin of the IRS13E group of stars is. It is conceivable that such a cluster could have been formed in an accretion disk as described by  Milosavljevic \\& Loeb (2004) and Nayakshin et al. (2005). However, numerical simulations have shown that it is difficult to explain the formation of such a dominant feature in a star-forming disk (Nayakshin et al. 2007). To support the cluster in-fall scenario, an intermediate-mass black hole (IMBH) was proposed to reside in the center of the cluster (Maillard et al. 2004). The existence of an IMBH results in a higher dynamical friction and a higher stability of the in-falling cluster such that this process in general is more efficient. It would allow compact clusters to reach the central parsec of the Galaxy (with $>$10$^{6}$\\solm)  within their lifetimes (Hansen \\& Milosavljevic 2003, Berukoff et al. 2006, Portegies Zwart et al. 2006, but see Kim et al. 2004 for a characterization of the problems with this hypothesis). However, the presence of the IMBH that would be required to stabilize IRS13E is highly disputed. Sch\\\"odel et al. (2005) analyzed the velocity dispersion of cluster stars. The authors found that the mass of such an object should be $\\ge$7000$\\,$\\solm. However, both the X-ray (Baganoff et al. 2003) and the radio (Zhao \\& Goss 1998) source at the position of IRS13E can be explained by colliding winds of high-mass-losing stars (Coker et al. 2002, Zhao \\& Goss 1998). These findings make the presence of an unusual and massive and dark object in IRS13E unnecessary.  In addition to the IRS13E association there are a number of faint infrared excess sources that may be essential for the discussion of star formation in the central stellar cluster (see Eckart et al. 2006a and Perger et al. 2008). In some cases there is spectroscopic evidence for them being associated with young stars (Perger et al. 2008). Here we present new proper motion and new spectroscopy data on these objects \\footnote{Based on observations collected at the European Southern Observatory, Chile}. These comprise first K$_s$-band proper motion measurements of sources that include the IRS13N association, and several sources at a projected distance of only a few arcseconds from SgrA*.  Approximately 0.5'' north of IRS13E, a small cluster of unusually red compact sources has been reported (IRS13N; Eckart et al. 2004). Muzic et al. (2008) gave a detailed analysis of this cluster and showed that an orbital analysis supports that the cluster is a comoving group of young stars. A strong infrared excess is due to the emission of warm dust (T $\\sim$$\\,1000$$\\,$K; Moultaka et al. 2005). The authors proposed two possible explanations for the nature of IRS13N: (1) objects older than a few Myr and similar to bow-shock sources reported by Tanner et al. (2005), or (2) very young stars (0.1 - 1 Myr old). The latter scenario implies more recent star formation than has been assumed so far for the GC environment.  Recently, Gillessen et al. (2012a) reported a fast-moving strong infrared excess source, which they interpreted as a core-less gas- and dust cloud approaching SgrA*. Owing to its location and because it is apparently on an elliptical orbit around SgrA*, we refer to it in the following as the dusty S-cluster object (DSO). The community has started to call the newly found fast-moving object G2 (e.g.  Burkert et al. 2012), reminiscent of the gas cloud G1 found by Clenet et al. (2005). As we show in the present paper, it cannot be excluded that the newly found object is instead a star with a gas/dust shell or disk around it, and not a core-less gas cloud as claimed by Gillessen et al. (2012a). We therefore propose to call it DSO rather than G2 since the acronym DSO (Dusty S-cluster object) reflects the true nature of the source in a more realistic way. While essential properties of this source have already been described by Gillessen et al. (2012a), no continuum identification of it shortward of 3.5$\\mu$m has been published to date. In other cases of near-infrared excess sources (e.g. see Fig.~14 in Eckart et al. 2006a) a detailed investigation of their nature had not yet been carried out. We also present first K$_s$-band proper motion studies of the DSO.  After the introduction we describe the observations and data reduction in section~\\ref{reduction}. In the results section~\\ref{results} we derive the proper motions and then outline the physical properties of the sources in section \\ref{discussion}. In section~\\ref{summary} we give a summary and the appendix contains additional maps and graphs.   \\begin{figure*}[ht!] \\centering \\includegraphics[width=19cm,angle=-00]{eckartfig-01.eps} \\caption{\\small Proper motion vectors of dusty objects in the central stellar cluster of the Milky Way (see Tab.~\\ref{Tab:ExcessPM}). We show the May 10, 2003, adaptive optics $L'$-band image of the central $4\"\\times4\"$. The arrows indicate the uncertainties in amount (red or yellow portion of the otherwise green arrow) and angle of the measured proper motion derived from data between 2002 and 2008 in the H, K$_s$, and L'-bands (see Tab.~\\ref{Tab:useddata}). } \\label{eckartfig-01} \\end{figure*} ",
        "conclusions": "\\label{discussion}  \\subsection{Identification and proper motions of IRS13N sources} The derived stellar 1D velocity dispersion at that location is three times higher than the root mean square difference between the proper motions of the sources identified in the K$_s$ and $L'$-band. In addition, the difference between our 2005.83 1D IRS13N source positions in the K$_s$-band and the $L'$-band positions by Muzic et al. (2008) at a mean epoch closer to 2005 is about 20~mas, i.e., two thirds of a K$_s$-band pixel of 27~mas size. This is on the same order (or below) of what is expected as the maximum change in position due to proper motions of the sources during the time difference between these mean epochs. We conclude that over the past 10 epochs we used to analyze the motion of the IRS13N sources the K$_s$-band sources can indeed be identified as being the co-moving counterparts of the $L'$-band sources. This also means that the $L'$-band sources are not unrelated (core-less) dust clouds as speculated by Fritz et al. (2010).  It is likely that not all sources $\\alpha$ to $\\kappa$ that had been discussed in Eckart et al. (2004) belong to the same association. We pointed out already in Eckart et al. (2004) that $\\kappa$ has significantly bluer colors than the other IRS13N objects. Infdeed, it has not yet been identified with an $L'$-band source. Our K$_s$-band proper motion analysis also shows that it is moving south and west rather than predominantly north as all the other IRS13N members. From this we conclude that $\\kappa$ does not belong to the small IRS13N cluster of co-moving infrared excess sources.  Source $\\alpha$ is farthest south of the Eckart et al. (2004) sources and has also a velocity offset from the other IRS13N sources. Here we give the $L'$-band velocities derived by Muzic et al. (2008) a higher weight, since in K$_s$-band the source is apparently located in a crowded region as well as in the wings of the bright IRS13E member E1. We conclude that $\\alpha$ is also probably not a member of the small IRS13N cluster of co-moving sources. We are therefore left with the sources $\\beta$ to $\\eta$ that have been identified as co-moving and whose dynamics were analyzed in detail by Muzic et al. (2008).  Source $\\eta$ is significantly brighter in $H$-band than all other IRS13N sources. It also has a the highest velocity discrepancy between K$_s$ and L' with a difference of $\\sim$150 km/s. If we assume that the IRS13N sources are embedded in their own dust disk or envelopes and that they all have a random orientation towards the observer, this implies an average opening angle of 180$^o$/6=30$^o$ under which the source can be seen unobscured. From this one may conclude that the circum-stellar material almost entirely surrounds the objects. If the sources have disks, it means either that not all material has settled into the disk, or that these disks are as thick as they are long in diameter, or that they are significantly warped. All of this could be expected if the sources indeed belong to a dynamically young compact cluster of co-moving young stellar objects.  The velocity dispersions of sources $\\beta$ to $\\eta$ corrected for their measurements uncertainties is about 45 km/s for the L'-band data discussed by Muzic et al. (2008). For the less well determined motions derived from our K$_s$-band identifications this value is about a factor of 1.5 higher. The implications for the possibly associated stellar mass of the IRS13N association has been discussed in detail by Muzic et al. (2008). The result is that if the IRS13N cluster is gravitationally bound, these velocity dispersions imply a stellar mass that should be detectable. Hence, it is likely to be unstable, implying that the association and its members are (dynamically) young.  \\noindent \\begin{figure} \\centering \\includegraphics[width=8cm,angle=-00]{eckartfig-35.eps} \\caption{\\small Distinguishing the DSO against stars S57, S23, S54, and S63 in 2008 in L'- and K$_s$-band. Only an upper limit can be given in the H-band. } \\label{eckartfig-35} \\end{figure}  \\noindent \\begin{figure}[ht!] \\centering \\includegraphics[width=9cm,angle=-00]{eckartfig-13.eps} \\caption{\\small Positions of the DSO relative to S2 and SgrA* plotted in comparison to the orbital tracks provided by Gillessen et al. (2012a) (red for the DSO, blue for S2, SgrA* is at the position of the black dot at the origin). The results for the K$_s$- and L'-band are shown in panels a) and b) (see data in Tab.~\\ref{Tab:DSO}). Thin lines indicate the approximate expected positions of the DSO on its orbit from 2002.0 to 2012.0 (based on our L'-band reductions and Gillessen et al. 2012a). The epochs are color-coded and the measurement uncertainty is shown as a black cross (see \\ref{eckartfig-24} in the appendix). } \\label{eckartfig-13} \\end{figure} \\subsection{Nature of the IRS13N sources} \\label{resultsIRS13N}  In this paper we used the SINFONI line and continuum images for a qualitative interpretation of the nature of the extended and stellar sources only. A quantitative analysis of the line fluxes is beyond the scope of the paper and will be published elsewhere. The comparison between the Br$\\gamma$ line emission and the near-infrared continuum emission of the stars in the IRS13E/IRS13N field in Fig.~\\ref{eckartfig-07} shows that the line emission peaks on one of the brightest components (E2) of IRS13E, covers the entire IRS13E cluster and extends to the IRS13N complex. The strength of the emission is correlated with the presence of stars in the field and is clearly not only associated with the stars, but comprises a diffuse component as well. The maps of the [FeIII]-emission depicted in Fig.~\\ref{eckartfig-08} show that the line flux is peaked both on IRS13E (E2) and IRS13N. The [FeIII] flux also shows contributions from a diffuse component. In IRS13N the [FeIII] integrated line flux density peaks to within $\\le$0.1'' close to sources $\\epsilon$ and $\\delta$. Zhao et al. (2009) found a radio source labeled K23 that lies between $\\epsilon$ and $\\delta$ and has to within 3-4$\\sigma$ the same velocity as the mean velocity of the NIR counterparts of $\\epsilon$ and $\\delta$.  The integrated [FeIII] line emission is also peaked on a source about 0.15'' north of $\\eta$ with offsets from SgrA* of $-3.25``$$\\pm$0.03`` in right ascension and $-0.45``$$\\pm$0.03`` in and declination. This position agrees with that of source K21 identified by Zhao et al. (2009) to within about 0.1''. The authors found a proper motion velocity of $v_{\\alpha}$=-224$\\pm$16~km~s$^{-1}$ and $v_{\\delta}$=-421$\\pm$20~km~s$^{-1}$ toward the NNW. In the $L'$-band we do not find a clearly defined point source at the position of K21. The continuum  flux density extension from IRS13N at that position is consistent with a source that is about a magnitude fainter (i.e. $L'$$\\sim$12) than source $\\eta$. At K$_s$-band no clear association to a source brighter than K$\\sim$16.3 can be established. This indicates that K21 is either associated with a stellar source even more reddened than the IRS13N sources or that the source K21 is a core-less dust source bright in [FeIII]-emission. The integrated line emission of Br$\\gamma$ is also extended and covers the position of K21. A spectrum of the infrared flux density at the position of K21 is shown in Fig.~\\ref{eckartfig-11}.  The comparison between the CO absorption and the near-infrared continuum emission in Fig.~\\ref{eckartfig-07} shows that the IRS13N complex is clearly located at a minimum of the CO-absorption in the IRS13E/IRS13N field. The absorption is concentrated on a few stars in the field to objects about 0.5'' east, west, and north of IRS13N. The minimum  in the absorption strength centered on IRS13N comes close to the background value due to the overall central stellar cluster. Given the contribution through PSF wings of these stars at the location of IRS13N, this is fully consistent with the statement that most of the IRS13N sources (about 0.5'' north of IRS13E) do not show any significant CO line absorption.   The comparison of the radio data by Zhao et al. (2009) (Fig.~\\ref{eckartfig-09}) shows that the radio continuum emission peaks to within $\\le$0.10'' at the peaks or contour line excursions of the [FeIII] line emission. The radio continuum is more peaked on the IRS13E source E1. The correspondence to the [FeIII]-emission (in particular the 2.242$\\mu$m line emission) is best. The Br$\\gamma$ emission is clearly more extended than the bright radio continuum emission, highlighting that the radio emission is more dominated by the compact stellar sources than the more diffuse emission traced by Br$\\gamma$.  The combined spectra on IRS13N and the two off-positions east and west (Figs.~\\ref{eckartfig-03} and \\ref{eckartfig-10}) demonstrate that the emission in the Br$\\gamma$, H$_2$ and in particular in the [FeIII] lines is brighter on IRS13E than in the off-positions. The spectra also demonstrate that the integrated CO-band head absorption toward IRS13E is as strong as toward the off-positions. This indicates that the CO-absorption at this position can be fully explained by the CO-absorption in the PSF wings of the bright nearby late-type stars and the overall stellar cluster background.   \\subsection{NIR colors and spectra of the IRS13N cluster} \\label{COhead}  The analysis of the HK$_s$L' (i.e. a significant section of the SED) colors presented in Fig.16 shows that they - if dereddened - would not come to lie even close to the positions expected for colors of pure dust-emitting sources but rather lie at the locations expected for objects that show a mixture of stellar and dust emission (see also Eckart et al. 2004).  Compared to the immediate surroundings (except for the IRS13E cluster), the sources in the IRS13N region are clearly correlated with a minimum in CO absorption and an excess in Fe- and H-recombination line emission. While the excess in emission may in part be due to a local concentration of gas associated with the 13N association or possibly with the mini-spiral, the lack of CO absorption can be attributed to the predominantly early spectral type, as suggested from the NIR colors of the sources (Eckart et al. 2004).  While NIR CO absorption and emission is often present, it is not a necessary feature to be observed toward young massive stars. Out of 201 young massive stars in M17, Hoffmeister et al. (2006) found 30\\% to have featureless spectra. Half of those show an infrared excess that corresponds to basically the same percentage of the remaining 60\\% of the sources that show CO in absorption or emission. Almost a third of the featureless sources also tend to show X-ray emission and a fifth shows simultaneous infrared excess and X-ray emission.  Featureless sources have the tendency to be brighter in the NIR compared to other young stellar objects. From the distribution in Fig.~3 in Hoffmeister et al. (2006) we can conclude that featureless objects are typically up to two magnitudes more luminous compared to CO emission or absorption sources. They are apparently later than spectral class B3.     Compared to previous CO bandhead studies in star-forming regions, the investigation by Hoffmeister et al. (2006) is the most extensive and a fairly unbiased one. Casali \\& Eiroa (1996) argued that Class\\,II sources tend to show CO absorption, while Class\\,I sources are featureless. The thermal continuum emission from dusty in-falling envelopes can exceed the combined stellar and disk photospheric emission by almost an order of magnitude (Calvet et al. 1997, Hoffmeister et al. 2006). Hence, this process can result in featureless spectra. In addition, the veiling of an envelope will also weaken the CO absorption lines. The presence of envelopes with large-scale heights is consistent with the presence of a single bright source $\\eta$ in addition to five faint objects, as elaborated above. The presence of a disk, on the other hand, is likely to amplify the CO feature. Therefore, Hoffmeister et al. (2006) concluded that the 60 featureless objects in their sample of 201 sources in M17 are very likely Class\\,I sources that have just started to build up an infrared excess that is progressively moving toward the MIR and NIR domain.  From model calculations Hoffmeister et al. (2006) infered that the observed CO absorption is most likely a sign of heavily accreting proto-stars. The mass accretion rates may be above $10^{-5}\\,$M$_\\odot$\\,yr$^{-1}$. High accretion can be provided through angular momentum loss within a disk. We conclude that the absence of that feature may indicate that these stars are currently not strongly accreting or have not yet reached that phase. This would be the case if the luminosity of the in-falling envelope is strong against the emission from the disk.  We infer from the combination of the fact that the stellar sources are in a co-moving group (Muzic et al. 2008, Eckart et al. 2004) and the proper motions and the spectroscopy presented here (see below) that the IRS13N sources may be featureless young stellar objects, like Class\\,I sources, in which a bright NIR disk has not yet been formed but in which the NIR/MIR excess is still dominated from surrounding and in-falling material.  \\subsection{Age of the IRS13N sources} \\label{IRS13Nages}  Possible stellar identifications of the IRS13N sources have already been discussed extensively by Eckart et al. (2004) and Muzic et al. (2008). The fact that our data supports the finding that the sources ($\\beta$ to $\\eta$; see below) are indeed co-moving poses serious problems to source identifications that involve stars with ages $>$3~Myr. Here we highlight a few more aspects concerning the identification of the sources with young stars:  Broad-band colors can give information on the nature of sources. In Fig.~\\ref{eckartfig-12}  (left) we compare the colors of dusty objects in the central few arcseconds of the Milky Way to the colors of young stellar objects as given by Hoffmeister et al. (2006). The sources of IRS13N are more highly extincted than the remaining dusty objects in the field. Their dereddened colors agree with those young dust enshrouded objects (details in Eckart et al. 2004). The assumption (see section~\\ref{COhead}) that the IRS13N cluster members are Class~I sources implies that they are young. In a recent study of stellar and circumstellar properties of six Class I proto-stars Prato et al. (2009) estimated their ages as $<$2~Myr. However, even the ages of young Class~II/Class~III sources are estimated to range between 0.1 and 1~Myr (Green \\& Meyer 1995, Luhman \\& Rieke 1999, see also van Kempen et al. 2009). This is consistent with the assumption of Eckart et al. (2004) that the IRS13N sources are good candidates for being YSOs or young Herbig Ae/Be stars with ages that may range between 0.1 and about $\\sim$3~Myr. Molinari et al. (2008) found that the time scales of the formation of stars are about 1-4$\\times$10$^5$ years, with shorter times for higher masses (6 to 40 \\solm) (Hillenbrand et al. 1992, Fuente  et al. 2002, Ishii et al. 1998). The orbital time scale for the combined motion of the entire IRS~13E/IRS13N complex and the mini-spiral gas of about 10$^{4}$ years (see also Muzic et al 2008) is short compared to the plausible ages of the massive young stars. We also assume that their associated envelopes and disks can survive the hostile GC environment for at least 10$^4$ years (see discussion in Scally \\& Clarke 2001).      \\begin{table*}[htb] \\begin{center} \\begin{tabular}{rrrrrrrrr}\\hline epoch & DSO &  DSO & DSO & DSO & S2 & S2\\\\ & K$_s$ &  K$_s$& L' & L' & HK$_s$L' & HK$_s$L'\\\\ & $\\Delta$$\\alpha$ & $\\Delta$$\\delta$ & $\\Delta$$\\alpha$ & $\\Delta$$\\delta$ & $\\Delta$$\\alpha$ & $\\Delta$$\\delta$  \\\\ \\hline 2002.42  &       &           & 0.255 &-0.130 &   0.020 &    0.000 \\\\ 2003.54  &       &           & 0.249 &-0.097 &   0.041 &    0.079 \\\\ 2004.66  &       &           & 0.255 &-0.121 &   0.034 &    0.125 \\\\ 2005.46  &       &           &       &       &   0.027 &    0.147 \\\\ 2005.83  &       &           & 0.232 &-0.095 &   0.024 &    0.150 \\\\ 2006.55  & 0.213 &   -0.094  & 0.225 &-0.082 &   0.011 &    0.168 \\\\ 2007.33  & 0.196 &   -0.074  & 0.208 &-0.083 &   0.000 &    0.176 \\\\ 2008.49  &       &           & 0.186 &-0.073 &  -0.015 &    0.179 \\\\ 2008.56  & 0.187 &   -0.065  &       &       &  -0.016 &    0.179 \\\\ 2009.47  & 0.165 &   -0.055  & 0.182 &-0.048 &  -0.026 &    0.177 \\\\ 2010.53  & 0.147 &   -0.033  &       &       &  -0.035 &    0.173 \\\\ 2011.55  & 0.120 &   -0.022  & 0.144 &-0.040 &  -0.044 &    0.167 \\\\ 2012.54  & 0.101 &   -0.017  & 0.108 &-0.023 &  -0.053 &    0.161 \\\\ \\hline \\end{tabular} \\end{center} \\caption{K$_s$-, and L'-band positions of the DSO and the star S2 relative to the SgrA* position. Typical uncertainties range between 13~mas and 25~mas for the DSO and (see Fig.~\\ref{eckartfig-24}) less than 10~mas for the position of S2. } \\label{Tab:DSO} \\end{table*}    \\begin{table*}[htb] \\begin{center} \\begin{tabular}{lrrrrrrrr}\\hline name & flux & epoch & $\\Delta$$\\alpha$ & $\\Delta$$\\delta$ & pm$_{\\alpha}$~~~ & pm$_{\\delta}$~~~ & acc$_{\\alpha}$~~~ & acc$_{\\delta}$~~~ \\\\ & ratio & year & mas & mas & mas yr$^{-1}$ & mas yr$^{-1}$ & mas yr$^{-2}$ & mas yr$^{-2}$ \\\\ \\hline S23 & 1.01 & 2005.47&  307.4&   -89.1&  14.81&   11.17&  -0.953&   0.525\\\\ S57 & 1.24 & 2007.46&  393.6&  -147.4&   9.99&    4.05&   0.000&   0.000\\\\ S63 & 0.83 & 2004.00&  245.0&  -150.0&  13.00&   -3.60&   0.000&   0.000\\\\ DSO & 0.30 & 2002.00&  295.0&  -160.0&  17.60&  -14.00&  -0.160&  -0.240\\\\ S54 & 1.36 & 2006.59&  115.4&   -60.1&   1.05&   26.90&   0.000&   0.000\\\\ \\hline \\end{tabular} \\end{center} \\caption{Epochs, flux ratios, positions, proper motions, and accelerations used to simulate the DSO and the neighboring stars with results shown in Figs.~\\ref{eckartfig-14-1}, \\ref{eckartfig-14-2}, and \\ref{eckartfig-14-3}. The flux ratios have been calculated with respect to the mean flux derived (m$_K$=17.6) from the K$_s$-band magnitudes of S23, S54, S57, and S63 as given by Gillessen et al. (1999). They correspond (with an estimated uncertainty of 30\\%) to the relative fluxes suggested by the modeling presented in the figures listed above. From the quality of the fit in these figures we estimate that the uncertainties are 10 mas for the coordinates, 1 mas~yr$^{-1}$ for the proper motions, and  0.1 mas~yr$^{-1}$ for the accelerations. } \\label{Tab:DSOsim} \\end{table*}  \\noindent \\begin{figure} \\centering \\includegraphics[width=9cm,angle=-00]{eckartfig-33.eps} \\caption{\\small Decomposition of the DSO spectrum including our K$_s$-band detection and H-band limit. Here we demonstrate that a mixture of dust and stellar contribution is possible. The points correspond to the L- and  K$_S$-band magnitudes, and H-band upper limit from this work, and the M-band measurement and K-band upper limit of Gillessen et al. (2002). Red and blue dashed curves also show their 550\\,K and 650\\,K warm dust fits. In solid blue, red, and magenta lines the emission from three different possible stellar types of the DSO core are plotted. Any of these stars embedded in 450\\,K dust (solid brown line) can produce the black line that fits all the NIR DSO photometric measurements. } \\label{eckartfig-33} \\end{figure}   \\noindent \\begin{figure*}[ht!] \\centering \\includegraphics[width=19cm,angle=-00]{eckartfig-12.eps} \\caption{\\small Left: Colors of dusty objects in the central parsec of the Milky Way compared to the colors of young stellar objects as given by Hoffmeister et al. (2006). The labels COAS and COES indicate that the sources show CO absorption  or emission in their spectra. Right: Colors of dusty objects in the central parsec of the Milky Way compared to colors of pure dust (temperatures given in red numbers) and mixtures of a stellar and a dust contribution (fractional numbers given in black; see also Glass \\&  Moorwood 1985). Here we demonstrate that a mixture of dust and stellar contribution is possible. Because the H-K$_S$ color of DSO is an upper limit, its NIR emission can be explained as pure dust at $550\\,{\\rm K}$ reddened by $\\sim 27\\,{\\rm mag}$, or as a star embedded in warm dust, e.g., if its color were detected to be ${\\rm  H-K_S\\sim 3.5\\, mag}$ (dashed circle), after dereddening by 27\\,mag, its emission would be consistent with 20\\%stellar and 80\\% dust contributions. } \\label{eckartfig-12} \\end{figure*}   \\noindent \\begin{figure*} \\centering \\includegraphics[width=15cm,angle=-00]{eckartfig-28.eps} \\caption{\\small Comparison of DSO images in the a) Ks- (13.2~mas pixel scale) and b) L'-band (27~mas pixel scale). The image section size is 1''$\\times$1''. In c)  and d) a Gaussian of 98~mas FWHM and a PSF from the image have been scaled to the DSO peak brightness and subtracted at its position. An astrometric grid of red circles marks some bright stars that can be used as positional references to compare the images at different wavelengths. } \\label{eckartfig-28} \\end{figure*}   \\subsection{Broad-band spectrum of IRS13N} Inspection of the 1.2'' angular resolution 1.3mm map of the mini-spiral presented by Kunneriath et al. (2011) shows that the IRS13N complex is associated with 1.3mm continuum emission. Just north of IRS13E it shows itself as a compact source component with a flux density of 20$\\pm$3~mJy (see Fig.~\\ref{eckartfig-16}). Arithmetically de-convolved with the beam, its size is less than 0.6~arcsec. In NE direction it is partially blended with IRS13E and a northern neighboring mm-source. In the 13mm map presented by Zhao \\& Goss (1998; Fig.3/PLATES therein) we can read off the map in their Fig.~3 an extended flux density of 11$\\pm$2mJy with a compact source of about 3mJy. These flux densities are consistent with a dominant nonthermal stellar wind contribution with an approximate mass load of 10$^{-5}$ \\solm~yr$^{-1}$ (e.g. Montes et al. 2009; Lang et al. 2005). This a mass load is typical for young luminous stars (Drake \\& Linsky 1989, Kudritzki et al. 1999, Crowther 2001).  Eckart et al. (2004) quoted H, K$_s$ and $L'$ flux densities for individual IRS13N sources. Viehmann et al. (2005, 2006) gave NIR/MIR flux density values obtained in an $\\approx$1'' diameter circular aperture on IRS13N. In the J- and $H$-band and probably even in the K$_s$-band these 1'' aperture flux densities are dominated by confusing fore- and background sources of the central cluster and the values by Viehmann et al. (2006) can probably only be taken as an upper limit. Due to the very red colors of the sources the 1'' flux densities at longer wavelengths are most likely characteristic for the IRS13N sources. However, at these wavelengths dust emission becomes dominant. Our investigation of the K$_s$- and $L'$-band proper motions shows that the dust-emission-dominated $L'$-band flux densities can be fully associated with the stellar sources found at shorter wavelength. Hence, we assume that a dominant portion of the emission longward of the $K_s$-band can in fact be attributed to the IRS13N sources.   \\subsection{Proper motions and nature of the central dusty sources}  Examination of the proper motions of the dusty sources in the central few arcseconds shows that 5 of the 10 L'-band excess sources identified in Eckart et al. (2004, see Fig.~\\ref{eckartfig-01}) have K$_s$-band counterparts. The color-color diagram in Fig.~\\ref{eckartfig-12} (left) shows that the colors of dusty objects in the central few arcseconds of the Milky Way are bluer than those of the IRS13N sources. Correcting for the Galactic foreground extinction of approximately 27 magnitudes in the visible places them at the location of the young stars discussed by Hoffmeister et al. (2006). The IRS13N sources are more highly extincted than the remaining dusty objects in the field. Correcting for the Galactic foreground extinction places them at the location of extincted young stellar objects in Fig~\\ref{eckartfig-12}i (left). Sources $\\beta$ to $\\eta$ in IRS13N are therefore excellent candidates for being young stars with an average extinction around $A_V \\sim 27$.  \\subsection{Pure dust versus photospheric emission} \\label{pure} Here we investigate if the NIR emission of the infrared excess sources can be caused by pure dust sources or if significant contributions from photospherice emission of stars need to be included. If the objects were pure dust sources their observed colors must be consistent with reddened colors of pure dust sources. For the IRS13N this is apparently not the case. The DSO is clearly detected in the K$_s$-band at m$_K$=18.9. Using A$_K$=2.22 and approximating $L_K$[\\solar] with $1.1 \\times 10^4 \\times D[Mpc]^2 \\times S_K [mJy]$ that results into $S_K=0.13 mJy$ and $L_K$ [\\solar] =92$\\times$10$^{-3}$\\solar. This is 46 times more than the luminosity in the Br$\\gamma$ line reported for the head of the DSO by Gillessen et al. (2012b). This implies that the 2$\\mu$m luminosity is dominated by the continuum rather than the Br$\\gamma$ line emission. In the H-band we find an upper limit m$_H$$>$21.2 if we consider the background level from the neighboxuring stars S23, S54i, and S63 in 2008, and the 3$\\sigma$ variation at the position of the DSO in the K$_s$- and L'-band. With the DSO L'-band brightness of m$_{L'}$=14.4$\\pm$0.3 obtained from the observing epochs listed in Tab.~\\ref{Tab:useddata} we obtain H-K$_s$=2.3 as a lower limit (Fig.~\\ref{eckartfig-12} right). This results in a measured color K$_s$-L'=4.5. With A$_H$=4.21, A$_{K_s}$=2.22 and A$_{L'}$=1.09, A$_{M}$$\\sim$A$_{L'}$ (e.g. Viehmann et al. 2005, Lutz et al. 1996) the absolute magnitudes can be found and  plotted in a spectrum (Fig.~\\ref{eckartfig-33}). In Fig.~\\ref{eckartfig-12} (right) we compare the data to the black-body line and extincted colors of different mixtures between a stellar and a dust contribution. (Glass \\&  Moorwood 1985). The pure 550~K dust colors (K$_s$-L'$\\sim$3.0) reddened by A$_V$=27 are consistent with the K$_s$-L' color of the DSO. However (with the current H-K$_s$ color limit and with potential H-K$_s$  colors that range up to the pure 550~K color), the dereddened positions of the DSO show that the DSO colors are also consistent with a mixture between a stellar and a dust contribution. Withn its current (H-K$_s$ color limited) position in the diagram, mixtures with 60\\% stellar emission and 40\\% dust with temperatures less than 550~K are possible. If the DSO will be detected half way between its current H-K$_s$ and that of pure 550~K dust (striped circle in Fig.~\\ref{eckartfig-12}), 20\\% stellar and 80\\% 550~K dust contribution are possible. In Fig.~\\ref{eckartfig-33} we show a possible spectral decomposition of the DSO using the M-band measurement by Gillessen et al. (2012a). In addition to their 550~K and 650~K fit and their m$_{K_s}$=17.8 limit (dashed lines) we plot our K$_s$-band  DSO detection and the H-band limit. The fit shows that the integrated flux can amount up to about 30~\\solar. We show a decomposition in which we assume that 50\\% of the current K$_s$ band flux is due to a late dwarf (blue line) or an AF giant or AGB star (magenta line). We used dust at a temperature of 450~K plotted in red (see dereddening of the DSO in the color color diagram in Fig.~\\ref{eckartfig-12} right). The sum of the spectra is shown by the thick black line (dashed at short wavelengths for the AF giant/AGB case).  Similarly, after dereddening the IRS13N sources for A$_V$=27$^m$, they miss the pure black-body line by about two magnitudes (see Fig.~\\ref{eckartfig-12} left). For fully dereddened objects an offset of two magnitudes in K$_s$-L' calls for temperatures of about 800~K or below with flux density contributions of a pure dust component of less than 50\\% (see also Fig.7 by Glass \\&  Moorwood 1985). This temperature estimate is consistent with the overall IRS13N temperature of around 650~K given by Fritz et al. (2010). This clearly implies that the IRS13N sources cannot be pure dust emitters but must have a stellar core. This is also fully consistent with the fact that the colors of these objects are similar to those of early dust enshrouded stars with colors shown the diagram in Fig.~\\ref{eckartfig-12} (left). Hence, for IRS13N sources the evidence for photospheric emission from a star is strongly supported by the H-band detections (Eckart et al. 2004 Tab.2 - here all sources except $\\epsilon$ and $\\gamma$ could be identified).  Inspection of Fig.~\\ref{eckartfig-12} (left), the black-body line shown in Fig.~\\ref{eckartfig-12} (right) and comparison to Fig.7 by Glass \\&  Moorwood (1985) shows that within the uncertainties the sources F1 and D7 are good candidates for being pure dust sources. Sources D6 and S90 will have mostly stellar colors after dereddening with a possible contribution of dust $>$1000~K. Sources F2, D2, D3, D4, and D5 are clear candidates for being objects that show a mixture from photospheric emission and emission from dust with temperatures well below 1000~K.   \\subsection{Statistical robustness of the DSO identifications} \\label{robust}  Here we analyze the statistical robustness of the DSO identifications in the H- and K$_s$-bands. In Sabha et al. (2012) and Eckart et al. (2012) we tabulated the probability for finding a blend star above the confusion limit in a single K$_s$-band spatial resolution element in the overall region of the S-star cluster and at the position of SgrA*. With two independent methods we determined the central K$_s$-band confusion limit observationally. Sabha et al. (2010, 2012) employed successive subtraction of stars in the S-star cluster. Witzel et al. (2012) used the shape of the SgrA* flare amplitude histogram. The otherwise straight power spectrum with which this histogram can be described turns over at a flare flux density at which a mere flare detection becomes a flare flux density measurement (details in Witzel et al. 2012). This indicates a confusion limit in the range of 0.5 to 1.5~mJy.  The result of the statistical investigation of the S-star cluster by Sabha et al. (2012) is that for a range of observationally supported parameters that describe the structure and luminosity function of the cluster, the authors can give the probability of finding a blend star in a single spatial resolution element above the confusion limit. Here we use a power law index of the projected spatial distribution of $\\alpha$=0.30$\\pm$0.05 and a K-band luminosity function exponent of KLF=0.18$\\pm$0.07. Projecting the S-star cluster onto the sky blend stars are created by line of sight overlaps of fainter stars within a finite angular resolution element. The probability of finding blend stars in the central 1.3$\\times$1.3 arcsec$^2$ overall can be as high as 1.0 and reaches a value of about 0.5 at the position of SgrA* itself. In the following we use these values although they can be much lower if the immediate vicinity of SgrA* is emptied of low-mass stars through stellar dynamical processes.  The probability of finding a K$_s$-band source in an L'-band resolution element (that corresponds to 2.8 K$_s$-band resolution elements) at a specific predicted position within the S-cluster is about 170 times lower than unity, i.e., 6$\\times$10$^{-3}$. As explained in Sabha et al. (2012), the live time of these blend stars (also indicated through observations) is about 3 years during which they dissolve again into individual objects with flux densities below the confusion limit due to their proper motions. If one aims to explain the 2007-2012 K$_s$-band identifications of the DSO by blend stars, one requires about 2-3 of these transient apparent objects. This happens with a probability of 3$\\times$10$^{-7}$. Since there is no guarantee that apparent proper motions during the dissolving phase of the blend stars follow the observed orbit and since we have chosen upper limits to start with, these low-probability values can be regarded as a safe upper limits for a scenario in which the DSO can be explained as the result of blend stars. Even if a blend star is assumed, one can think of at least 5-10 proper motion velocity values and 5-10 directions that would be expected given the accuracy of 13-25~mas and the involvement of 2 to 6 epochs. This will lower the probabilities quoted above by a factor of 25-100. This implies that the combined probability of being confused by a comoving blend star is less than 10$^{-4}$.  However, the number of bright, single stars in the central region is high. We can therefore consider the likelihood of finding faint stars above the confusion limit at the position of the DSO over a time span of several years. From the KLF analysis by Sabha et al. (2010, 2012) we infer a maximum of about 40 stars fainter than 16 magnitudes (dereddened) within the central 1.3''x1.3''. This implies 0.2 sources per resolution element and observing epoch. For six consecutive epochs this results in a probability of about a percent to find a random configuration of the central cluster such that a faint star is found at consecutive positions given by the L'-band detections. If we include the immediate neighborhood (i.e., 9 to 25 surrounding resolution elements), the probability can rise to 25\\%. With a limited accuracy of determining the direction and value of the proper motion, this implies that the likelihood of finding a star that is close in position and velocity is in the percentage range. This may explain the confusion of the DSO with S63 in the interval 2002 to 2007.  In summary, the contamination with a bright field star over the entire time interval 2002 to 2012 is likely (around a few percent)  but being confused by blend stars is very unlikely ($\\le$10$^{-4}$). For the IRS13N and all other infrared excess sources discussed here the projected stellar densities are even lower than those close to the center, such that serendipitous identifications are also not very likely.  \\subsection{Proper motions and identification of the DSO}  The analysis presented here shows that the dusty S-cluster object reported by Gillessen et al. (2012a) does indeed have a K$_s$-band counterpart. Its H-K$_s$-color limits indicate that it cannot be excluded that the DSO is an embedded star rather than a pure dust cloud. Gillessen et al. (2012a) showed that both in proper motions and radial velocities the trajectory of the source perfectly fits a closed elliptical orbit. The authors interpreted this source as a dust cloud that is approaching SgrA* and will be disrupted during the peri-bothron passage. The combination of the fragility of the cloud and the fact that it is on a bound elliptical orbit implies that the object must have been on that orbit for a couple of revolutions. In the harsh environment of the central stellar cluster the evaporation time scale for a pure dust cloud is incompatible with this fact. The dust cloud would disappear before it can establish a relaxed orbit. If the object is on a bound orbit as a cloud and since it will be heavily stretched out by tidal forces during the peri-bothron passages it should not be as compact as it currently appears. Variations in the shape should (with lower angular resolution than in the H- and K$_s$-band) be traceable in the bright L'-band continuum where it appears to be rather compact (Gillessen et al. 2012a). In the K$_s$-band the DSO is close to the confusion limit (see discussion in subsection \\ref{robust}). Given the strong and extended Br$\\gamma$ recombination line emission from the mini-spiral the situation may be similar for the DSO Br$\\gamma$ line emission. As in the case of the continuum identifications this may influence the brightness, position, shape, and velocity field of the source as seen in the light of the line emission (Gillessen et al. 2012a).  If the DSO is indeed a dusty star, then it may develop a bow-shock if it passes through an accretion wind from SgrA* quite similar to the sources X3 and X7 described by Muzic et al. (2010). Since the DSO as a dusty source is an obvious probe for strong winds possibly associated with SgrA*. Its presence and the fact that it has not yet developed a bow-shock, although it is already closer than X3 and X7, indicates that the wind from SgrA* is highly un-isotropic and possibly directed toward the mini-cavity (see Scenario IV below and Muzic et al. 2010 for a more detailed discussion).  Comparing the luminosity and the magnitude of the DSO (Tab.~\\ref{Tab:ExcessPM}) we obtained in the H- and K$_s$-band to the absolute magnitudes of OB and Be stars as presented by Wegner (2006) and stars of luminosity class I-V, we found that the DSO could preferentially be a K-dwarf (V) or an F-dwarf/subgiant (V-III) at no additional extinction intrinsic to the surrounding dust cloud or an A-dwarf/subgiant (V-III) at about 30 magnitudes of intrinsic visual extinction (see Fig.~\\ref{eckartfig-33}). The AGB or early evolutionary phase may be included here to support the dusty envelope. These identifications point toward a slightly later type as compared to the earlier-type B0-2.5~V stars in the S-cluster (Martins et al. 2008). This would imply a stellar mass of a few solar masses. However, the combination of local extinction (including geometrical effects like shadowing and volume-filling factors of the extincting material) and scattering may alter that identification and it could be even earlier or more luminous as well. This leads us to discuss four additional scenarios: \\\\ \\\\ {\\it Scenario I:} Since the DSO is prominent in its hydrogen line emission (Gillessen et al. 2012a) with a Br$\\gamma$ line with of 100~km/s in 2008 but rather faint in the H and K$_s$ continuum bands (see Tab.~\\ref{Tab:ExcessPM}), it may very well be a dusty, faint narrow-line Wolf-Rayet (WR) star. They are known to be less luminous than most other WR or OB giants, and there are a few of these objects in the GC field (see e.g. Moultaka et al. 2005, Paumard et al. 2006). The WR stars IRS7W with m$_K$=13.1, IRS7SE with m$_K$=13.0 (Martins et al. 2007) and WR2 with m$_K$=12.9 (Moultaka et al. 2005) are amongst the faintest known so far in the Galactic Center cluster (for faint WR stars in the Milky Way see also  Shara et al. 1999, 2012, Mauerhan, van Dyk, Morris 2011 with survey limits around m$_K$=15.0). Given these apparent magnitudes, the DSO would have to be an exceptionally faint WR star. The earliest WR stars with dense winds and strong free-free excesses may belong to the intrinsically faintest of the WRs (Mauerhan, Van Dyk, Morris 2011). As indicated by the increasing Br$\\gamma$ (Gillessen et al. 2012a), the harsh environment of the S-star cluster and the possibly repeated interaction with the SgrA* black hole may have altered the stars mass loss or shell size and hence its luminosity in the thermal infrared and the Br$\\gamma$ line width of an object that started out as a narrow-line WR star at larger separations from SgrA*. \\\\ \\\\ {\\it Scenario II:} Dong, Wang \\& Morris (2012) have carried out a multi-wavelength study of evolved massive stars in the GC. In their Fig.~12 they show a Pa$\\alpha$ equivalent line width (EW) vs. K$_s$-[3.6$\\mu$m] plot of these objects. From this plot it becomes obvious that for WC stars and the fainter WNL stars/OB super-giants the EW values are lower than 100$\\AA$ for K$_s$-[3.6$\\mu$m]$>$2 magnitudes and even drop significantly with further increasing infrared excess, reaching values of up to K$_s$-[3.6$\\mu$m]$\\sim$5. Dong, Wang \\& Morris (2012)  suggested that the free-free emission from the strong winds of the WNL stars/OB giants could potentially dominate the emission in the mid-infrared (Wright \\& Barlow 1975). This could explain the positive EW vs. K$_s$-[3.6$\\mu$m] correlation they found. These properties fit with the DSO well if we assume that K$_s$-[3.6$\\mu$m]$\\approx$K$_s$-L' and that a small EW in the Pa$\\alpha$ line is correlated with a small EW of the Br$\\gamma$ line. The Br$\\gamma$ line width in 2008 indicated in Fig.2 by Gillessen et al. (2012a) is on the order of or smaller than 100 km/s in 2008 and 200 km/s in 2011. Taking the Galactic foreground extinction into account, its position in the color-color diagram in Fig.~\\ref{eckartfig-12} (left) comes to lie in the WNL stars/OB supergiant domain with small EW and high infrared excess, as described by Dong, Wang \\& Morris (2012). \\\\ \\\\ {\\it Scenario III:} Murray-Clay \\& Loeb (2012) have proposed that the dusty S-cluster object is a proto-planetary disk that has been brought in from a young stellar ring. However, to change the ellipticity of the previously rather circular orbit in one single event to the high ellipticity orbit on which it is now requires a similarly violent interaction as the upcoming peri-bothron passage. If the object were only a core-less dust cloud, this event would therefore lead to an early destruction of the object at the time it enters the current orbit. If the source is indeed a core-less dust cloud, the only alternative is that the ellipticity of its orbit has been changed gradually. However, this implies many peri-bothron passages, imposes a conflict with evaporation timescales, and will most likely lead to destruction as well. The only alternative is indeed that the dusty S-cluster object has been formed on an orbit rather similar to the current one, or that it has been brought in gradually from outside the S-star cluster, as proposed earlier by Murray-Clay \\& Loeb (2012). This hypothesis is also supported by the analysis in the color-color diagram shown in Fig.~\\ref{eckartfig-12} (left). The authors show that there are two objects, D3 and the DSO, which are dominated by their dust emission in the $L'$-band, which would be the case if they were relatively young stellar objects. \\\\ \\\\ {\\it Scenario IV:} The DSO may be similar to the bow-shock sources X3 and X7 (Muzic et al. 2010) but has not yet developed an obvious bow-shock structure. For X3 and X7 Muzic et al. (2010) discussed possible stellar counterparts, including the possibility of late-type and main-sequence sources. However, the preference is clearly on late B-type main-sequence stars (B7-8V) and low-luminosity WR stars as well. Other explanations have problems: The central star of a planetary nebulae will remain at the required brightness for less than 1000 yr before it becomes too faint, and a main-sequence star cannot be an appropriate source of a dust-rich envelope, as is apparently observed for the DSO in the L -band.   \\subsection{Will the DSO be disrupted?} Gillessen et al. (2012a) have reported that the Br$\\gamma$ emission of the DSO is spread over 200~mas. Similar sizes are also predicted in the models presented by Burkert et al. (2012) and Schartmann et al. 2012. In Fig.\\ref{eckartfig-28} we compare images of the DSO in the Ks- and L'-band. We also show the L'-band images after removing the DSO using a Gaussian of 98~mas FWHM (i.e., the diffraction limit in the L'-band) and a star located 0.5'' west and 0.42'' north of SgrA* (i.e., a PSF obtained from an image section close to SgrA*; see Fig.~\\ref{eckartfig-01}). The analysis of the images shows that $>$90\\% of the DSO emission at 3.8$\\mu$m wavelength is compact (FWHM$\\le$20~mas). Only up to 10\\% of the flux density contained in the compact part of the DSO can be extended on the scale size of the PSF. This indicates that the warm ($\\sim$550~K) dust emission is very compact in comparison to the hotter ($\\sim$10$^4$~K) hydrogen recombination line emission. Alternatively, the MIR continuum emission contains a significant compact free-free contribution (see below). In this case, however, one would expect that the recombination line emission is dominated by contributions from this free-free component and should hence be more compact than observed.  Since the K$_s$-band identifications of the DSO suggest that it can also be associated with a star it is only its dusty envelope that can potentially be disrupted. To investigate that possibility we have to determine the location of the Lagrange point $L1$ through which mass can be transferred and compare it to the size of the dust shell or disk. In Fig.~\\ref{eckartfig-17} we show a sketch that demonstrates the location of the DSO and the motion of the Lagrange point L1 with respect to it. In panel a) in Fig.~\\ref{eckartfig-17} we show the situation about one year before the peri-bothron passage. The L1 point is rapidly approaching the DSO through its orbital motion toward SgrA*. At peri-bothron, $L1$ will also be closest to the star (panel b) in Fig.~\\ref{eckartfig-17}). If the star has a mass of about 1\\solm,  the separation of $L1$ from it will be about 0.1~AU. For a Herbig Ae/Be stars with 2-8\\solm ~that distance will be 0.2 and 0.5~AU. For a typical S-cluster stellar mass of $\\sim$20-30\\solm ~the separation will be closer to one AU. Interferometrically determined typical inner ring sizes for young Herbig Ae/Be and T~Tauri stars can be as small as 0.1-1~AU (Monnier \\& Millan-Gabet 2002), however, a disk or shell may be much larger than this. Based on its MIR-luminosity, Gillessen et al. (2012a) determined a source size of about 1~AU. Hence, during peri-bothron passage a significant amount of the dusty circumstellar material may pass beyond $L1$ and start to move into the Roche lobe associated with SgrA*.  \\noindent \\begin{figure} \\centering \\includegraphics[width=8cm,angle=-00]{eckartfig-17.eps} \\caption{\\small Sketch of the relative position and motion of the Lagrange point L1 and the DSO.} \\label{eckartfig-17} \\end{figure}    However, this  peri-bothron passage will only last for a few months, after which the DSO will move away from SgrA* at a speed of around a 1000 km/s within the first year. After one year the separation of the $L1$-point from the DSO is already about 1000 times larger than during the peri-bothron passage (panel c) in Fig.~\\ref{eckartfig-17}). This implies that any cold dusty material that passed the $L1$ toward the SgrA* will be overtaken again by the receding $L1$-point. However, the time will not be sufficient to completely disrupt the dust shell or disk of the DSO, especially if it is very compact. The material that passed he $L1$-point during the peri-bothron passage will not stay in the SgrA* Roche lobe and will not be immediately accreted. Only thin, hot material associated with a possible stellar wind of a hot and massive object may stay past the $L1$-point, but at the typical escape velocities of such volatile material of a few 100 km/s it would have left the star in any case.  Shcherbakov \\& Baganoff (2010) have discussed the feeding rate of SgrA* as a function of radius. Their modeling suggests that sources X3 and X7 (Muzic et al. 2010) are still in the regime in which most of the in-flowing mass is blown away again (see their Fig.3). This may partly explain the bow-shock structure of X3 and X7. During its peri-bothron passage the star S2 has been well within the zone in which matter of its stellar wind could have been accreted. However, no effect on the variability of SgrA* has been reported. The DSO peri-bothron will be at a larger radius than that of S2. This implies that if the stellar wind from its central star is not stronger than that of S2, no enhanced accretion effect will result from it. If the radius-dependent accretion flow is clumpy or un-isotropic, the DSO may still develop a bow-shock like X3 and X7 and a major portion of the cold and dusty material that was detached from the DSO during its peri-bothron passage will be blown away."
    },
    "1208/1208.3305_arXiv.txt": {
        "abstract": "{In May 1982, when Italy joined ESO, only two isolated neutron stars (INSs) had been identified in the optical: the Crab and Vela pulsars. Thanks to the ESO telescopes and the perseverance of a few Italian astronomers, now about 30 INSs have been identified in the optical/IR, and a new important channel in their multi-wavelength studies  has been opened. In this contribution, I review the major steps in 30 years of INS studies at ESO, highlight the role of Italian astronomers, and introduce future perspectives with the E-ELT. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0595_arXiv.txt": {
        "abstract": "The \\kepler\\ Mission is searching for Earth-size planets orbiting solar-like stars by simultaneously observing $>$160,000 stars to detect sequences of transit events in the photometric light curves. The Combined Differential Photometric Precision (CDPP) is the metric that defines the ease with which these weak terrestrial transit signatures can be detected. An understanding of CDPP is invaluable for evaluating the completeness of the \\kepler\\ survey and inferring the underlying planet population. This paper describes how the \\kepler\\ CDPP is calculated, and introduces tables of \\rms\\ CDPP on a per-target basis for 3-, 6-, and 12-hour transit durations, which are now available for all \\kepler\\ observations. Quarter~3 is the first typical set of observations at the nominal length and completeness for a quarter, from 2009 September 18 to 2009 December 16, and we examine the properties of the \\rms\\ CDPP distribution for this data set. Finally, we describe how to employ CDPP to calculate target completeness, an important use case. ",
        "introduction": "\\label{s:intro}  The \\kepler\\ Mission is a NASA Discovery mission designed to detect transiting extrasolar planets, performing near-continuous photometric observations of $>$160,000 carefully selected target stars in \\emph{Kepler's} 115 square degree field of view, as reviewed in \\citet{Borucki2010} and \\citet{Koch2010}. Scores of planets have been confirmed thus far{\\footnote{See \\url{http://kepler.nasa.gov/Mission/discoveries/}}, and three catalogues of planet candidates have been released: 705 candidates discovered in the first month of observations \\citep{Borucki2011a}, 1235 candidates discovered in the first 15 months of observations \\citep{Borucki2011b}, and 2321 candidates discovered in the first 18 months of observations \\citep{Batalha2012}. Although individual planetary systems continue to surprise and intrigue---see for example the Kepler-36 system \\citep{Carter2012}---we are now able to shift towards broader analysis of the underlying planetary populations \\citep{Borucki2011a,Howard2011,Youdin2011,Catanzarite2011} and trends, such as comparing single planet candidate systems to those with multiple planet candidates \\citep{Latham2011}, and the environments of small planet candidates compared to large planet candidates \\citep{Buchhave2012}.  The primary goal of the \\kepler\\ Mission is to ascertain the value of \\etaEarth, the frequency of Earth-size planets orbiting in the habitable zones of solar-like stars. Inferring the value of \\etaEarth\\ from the planet sample discovered by \\kepler\\ requires careful quantification of the detectability of each planetary candidate across the entire set of target stars. An essential aspect of measuring \\etaEarth\\ is accounting for the observation noise specific to each target star, and the subsequent impact on the detectability of the transit signature of the candidate, which is the topic of this paper.  \\kepler's transiting planet search (TPS) pipeline module \\citep{Jenkins2010,Tenenbaum2012}, which searches through the data for evidence of transit signatures, empirically determines the level of non-stationary noise for each light curve. This noise estimate is called the combined differential photometric precision (CDPP); it is a time series of the effective white noise as seen by a specific transit duration for each target star. To facilitate analysis and interpretation of the \\kepler\\ data, tables of the rms CDPP metrics on a per-target basis for transit durations of 3, 6, and 12 hours are provided online at the Mikulski Archive for Space Telescopes (MAST) website{\\footnote {\\tt http://http://archive.stsci.edu/kepler/}. Section~\\ref{sec:cdppcalculation} of this paper describes how CDPP is calculated by the \\kepler\\ pipeline. Section~\\ref{sec:tabledescription} provides a guide to the format and content of the tables. Section~\\ref{sec:q3cdpp} discusses the general characteristics of the Q3 rms CDPP values, and examines the distribution as a function of stellar type and position in the \\kepler\\ field. Section \\ref{sec:uses} describes some of the ways in which CDPP can be used for further analysis, and Section \\ref{sec:conclusion} contains the conclusions of the paper. ",
        "conclusions": "\\label{sec:conclusion}  We have presented here an introduction to the rms CDPP values being made available on a per-target, per-quarter basis at the MAST website for the \\kepler\\ planetary targets. For each set of data, these values provide a measure of the observed noise, which for each individual target translates to a limit on the detectability of transiting planets. These values are extremely important for the characterisation of the underlying planet population, since to first order, they provide on a star-by-star basis the observational noise level that sets the limiting planet signal detectable by Kepler."
    },
    "1208/1208.5009_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{introduction}  Indirect searches for Dark Matter, i.e. searches for `anomalous' features in cosmic rays (e.g. gamma-rays, neutrinos, positrons and anti-protons), have been proposed in the late 70's as a powerful way to reveal the existence of Dark Matter  annihilations in the Milky Way halo and beyond~\\cite{indirectgamma,indirectcharged}. These techniques are meant to give precious insights about the nature of the Dark Matter particle and its properties, assuming that a signal is seen. Yet there are several limiting factors which weaken their ability to elucidate the dark matter problem. In particular indirect detection requires a detailed knowledge of the astrophysical backgrounds and foregrounds and therefore depends on the present knowledge of astrophysical sources. To make a discovery one either has to carefully remove known (or modelled) background in order to expose the `anomalous' component or hope that the Dark Matter signal is well above the background and exhibits very clear features, which would be difficult to mimic by invoking astrophysical sources only.  \\medskip  Currently there are contradicting claims regarding whether indirect detection is giving clues of dark matter or not. On one hand, there are  possible anomaly detections which could be explained in terms of Dark Matter annihilation or decay. These include for example the positron excess, as seen in \\PAMELA\\ (and \\Fermi-LAT) data~\\cite{PAMELApositrons}, a possible feature in the $e^+ + e^-$ spectrum~\\cite{epem}, a claimed $\\gamma$-ray excess at $\\sim$ 10 GeV energies~\\cite{Hooper10GeV}  \\footnote{All these claims have possible drawbacks, cf \\cite{Boehm:2002yz,Boehm:2010kg,Bertone:2008xr,Cirelli:2009vg,Cirelli:2009bb,Cirelli:2009dv,Cirelli:2012ut}.} and, most recently, two possible $\\gamma$-ray lines at 111 and 129 GeV~\\cite{line130GeV,noline130GeV}. On the other hand, a large bulk of present astrophysical data essentially seem to validate the modelling of astrophysical background sources in the GeV-TeV range (disregarding these possible anomalies), and therefore enables one to set powerful constraints on the Dark Matter properties.  \\medskip By measuring the gamma-ray spectrum over a large energy range relevant for Dark Matter physics, the \\Fermi-LAT\\ collaboration has been able to  set stringent limits on the Dark Matter pair annihilation cross section into Standard Model particles. For example, using the diffuse $\\gamma$-ray emission in dwarf spheroidal (dSph) galaxies~\\cite{FermiDwarfs} and also in the Milky Way~\\cite{FermiDiffuseMW,Ackermann:2012qk}, the \\Fermi-LAT\\ collaboration has ruled out Dark Matter candidates with a total  annihilation cross section of $\\langle \\sigma v \\rangle = 3 \\times  10^{-26} \\ \\rm{cm^3/s}$ if $m_{\\rm DM} \\lesssim 30$ GeV. This constituted a remarkable milestone as such a value corresponds to that suggested by the thermal freeze-out scenario, which is generally considered as a strong argument in favour of Weakly Interacting Massive Particles (WIMPs).   \\medskip  These limits nevertheless weaken at higher DM masses, therefore allowing for heavier DM candidates with a larger pair annihilation cross section. For example, for $m_{\\rm{DM}} =$ 100 GeV the limit relaxes to $\\langle \\sigma v \\rangle \\lesssim 10^{-25} \\ \\rm{cm^3/s}$ while for $m_{\\rm DM} =$ 500 GeV, it reads $\\langle \\sigma v \\rangle \\lesssim 3 \\times 10^{-25} \\ \\rm{cm^3/s}$, which is one order of magnitude higher than the `thermal' cross section.  \\medskip DM models with such large values of the pair annihilation cross section  have actually been proposed over the last five years as a consequence of the excesses in $e^+$ and $e^++e^-$ fluxes. While they may remain hypothetical, discovering such a configuration would invalidate the WIMP `vanilla' model and either point towards the existence of non-thermal process in the Early Universe (possibly opening up an unexpected window on fundamental physics at high energies) or potentially call for more sophisticated mechanisms, such as Freeze-In and regeneration as proposed in \\cite{Hall:2009bx,Chu:2011be}. Explaining the observed dark matter relic density may remain nevertheless  challenging. For example, in \\cite{Williams:2012pz}, it was shown that candidates with a total annihilation cross section exceeding $\\langle \\sigma v \\rangle =  10^{-24} \\ \\rm{cm^3/s}$ (corresponding to a thermal relic density smaller than $3 \\%$) would be ruled out by the \\Fermi-LAT experiment if they were regenerated at 100$\\%$.  \\medskip  In addition to measurements of the $e^+$ and $e^+ + e^-$ spectra mentioned above, there is also the measurement of the galactic $\\bar{p}$ flux, presented by the \\PAMELA\\ collaboration~\\cite{PAMELApbar1,PAMELApbar2}. While extensive work was done to explain the electron/positron excesses in terms of Dark Matter annihilations (or decays), the implications of the absence of anomalies in the $\\bar{p}$ spectrum has remained relatively unexploited. Indeed only a relatively small number of works~\\cite{CKRS,salati,boehm,Evoli:2011id,Garny:2011ii, Asano:2011ik,Garny:2012eb} have dealt with it and shown that large Dark Matter annihilation cross sections can be constrained by the  \\PAMELA\\ data. Among the most interesting conclusions which have been reached let us cite for example that in \\cite{salati} constraints on the annihilation cross section into $b \\bar{b}$ were given (for the same mass range as is considered in this paper) and limits on the $W^+ W^- $ final state were mentioned for $m_{\\rm DM} = 1$ TeV and one specific set of propagation parameters. In \\cite{Asano:2011ik}, constraints on the $q\\bar{q}g$ were set for bino-like neutralinos.    \\medskip  The first aim of this paper is therefore to propose a more systematic analysis of these general anti-proton constraints on the DM annihilation cross section, including paying attention to the uncertainties associated with DM and astrophysical predictions. The  second aim of the present analysis is to demonstrate that these measurements can actually constrain the properties of specific DM scenarios, including the mass spectrum in the dark sector. To illustrate this, we will work within a `simplified' version of the phenomenological Minimal Supersymmetric Standard Model (pMSSM)~\\cite{Djouadi:1998di} in which all sfermion masses are set to 2 TeV, except for the stop and sbottom masses. The soft masses for the stop are allowed to be much lighter to obtain a Higgs at 125 GeV. In this scenario the only particles with masses below the TeV threshold are therefore the neutralino, chargino, the supersymmetric Higgses and the lightest stop and sbottom. Such a configuration of `light' gauginos and heavy sfermions may actually seem unnatural from a supersymmetric point of view (albeit close to split SUSY~\\cite{splitSuSy}) but it is supported by the unfruitful searches for squarks and gluinos at LHC, at least to some extent~\\footnote{Even though, admittedly, those negative searches may also be a sign that Supersymmetry is not realised at the TeV scale.}.  \\medskip With this very set up in mind, one can investigate scenarios where the neutralino pair annihilation cross section into $W^+ \\, W^-$  gauge bosons is enhanced (due in particular to the chargino exchange diagram). Such a large annihilation cross section gives both a significant anti-proton and diffuse gamma ray flux, together with a gamma ray line, and is therefore  potentially constrained by the \\PAMELA\\ and \\Fermi-LAT data. In Supersymmetry, such an enhancement is realised when the LSP neutralino is mass degenerated with the chargino, i.e. when the neutralino has a significant wino component. The combination of both \\Fermi-LAT and \\PAMELA\\ data is therefore expected to constrain the wino fraction of the lightest neutralino, thus realizing our second aim.  Note that constraints on the neutralino composition are also expected to be obtained in presence of a lower sfermion mass spectrum. However the effect of the chargino-neutralino mass degeneracy on $\\gamma-$ray and $\\bar{p}$ production would be much harder to characterise. Hence our choice in favour of a heavy sfermion mass spectrum.  \\medskip  The paper is organised as follows. In section \\ref{constraints} we derive generic constraints on the Dark Matter pair annihilation cross section into $W^+W^-$ from anti-proton data and recall the \\Fermi-LAT limits that are obtained from gamma-ray observations in the Milky Way and dwarf Spheroidal  galaxies. In Section \\ref{SuSy} we present the Supersymmetric model that we shall consider and explain how we perform the scans of the parameter space. Finally in Section \\ref{results} we apply the \\PAMELA\\ and \\Fermi-LAT  limits to our SUSY model and show that the anti-proton data can be more constraining than gamma-ray observations. We conclude in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  \\medskip In this paper we explicitly derived the constraints on the ${\\rm DM} {\\rm DM} \\rightarrow W^+ W^-$ annihilation cross section by using the \\PAMELA\\ anti-proton data and paying particular attention to the choice of propagation parameters and uncertainties on the astrophysical background. Our results are independent of the so-called \\PAMELA\\  positron excess and are obtained for two different (fixed vs marginalised) choices of the background spectrum; they are also consistent with the enhancement factor which was derived in \\cite{salati} and the detailed analysis of \\cite{Evoli:2011id}, for the cases where the propagation parameters overlap.  \\medskip We then compared these bounds with the most stringent gamma-ray limits which have been derived using the \\Fermi-LAT\\ measurements of the gamma-ray continuum spectrum from dSph galaxies, for the same DM annihilation channel and DM mass range. We found that the anti-proton constraints appear to be very competitive with the gamma-ray bounds. More precisely, choosing the `MED' propagation scheme, the $\\bar p$ constraints are slightly weaker than the $\\gamma$-ray ones when $m_{\\rm DM} \\lesssim  300$ GeV and slightly stronger for $m_{\\rm DM} \\gtrsim 300$ GeV.  On the other hand, the anti-proton constraints are stronger if we assume the 'MAX' set of propagation parameters and less powerful if we assume the `MIN' set. We also recall that the gamma ray limits themselves may be subject to some uncertainties related to the modelling of the DM profile in dSph  galaxies.  \\medskip Finally we applied as fiducial limits the $\\bar p$ constraints relative to `MED' and the marginalized astrophysical background to the neutralino LSP in a simplified version of the pMSSM, where we set all the sfermion masses (apart from that of the third generation) to the TeV scale. We found that the fiducial \\PAMELA\\ anti-proton and \\Fermi-LAT gamma-ray limits rule out small but non negligible neutralino-chargino mass splittings. In particular for $m_{\\chi_1^0} \\lesssim 150$ GeV, one can rule out mass splittings up to 20 GeV. Our results also suggest that pure wino or wino-like neutralinos are excluded if they are lighter than 450 GeV. Overall, this limit surpasses the bounds that can be set by using the XENON100 data and even in fact than the projected XENON1T limit.  \\medskip Hence from this work, we conclude that present indirect detection data already enable one to exclude regions of the parameter space where the neutralino-chargino mass splitting is small but non negligible. Since these regions are difficult to probe directly at the LHC, these findings show that \\Fermi-LAT and \\PAMELA\\ data constitute modern tools to explore the supersymmetric parameter space and even beat LHC (and also in fact Direct Detection) searches on their own territory, even though -- on the negative side -- they assume a regeneration of the relic density for neutralinos with a very large annihilation cross section."
    },
    "1208/1208.3255.txt": {
        "abstract": "{ We use the in-in or Schwinger-Keldysh formalism to explore the construction and interpretation of effective field theories for time-dependent systems evolving out of equilibrium.  Starting with a simple model consisting of a heavy and a light scalar field taken to be in their free vacuum states at a finite initial time, we study the effects from the heavy field on the dynamics of the light field by analyzing the equation of motion for the expectation value of the light background field.  New terms appear which cannot arise from a local action of an effective field theory in terms of the light field, though they disappear in the adiabatic limit.  We discuss the origins of these terms as well as their possible implications for time dependent situations such as inflation.} ",
        "introduction": "\\label{sec:Intro}  Effective theories provide a powerful method for understanding nature.  Their basic idea is to include only what is relevant for the particular phenomenon that is being described.  For example, to understand how a set of particles scatter and interact, it is not necessary to find a description that applies to all scales, but only one that works for those energies and momenta and precisions that can actually be measured.  All of the further details of nature at higher energies, or shorter distances, can be treated as though they act at a point from the perspective of what a lower energy experiment can resolve.  The symmetries, both of space-time and of the particles' interactions, are an important part of this picture.  In most particle physics applications, the set of symmetries is quite large, and quite constraining, and usually assumes an invariance under the Poincar\\'e transformations.  But other settings could have less symmetry.  In the very early universe, the space-time expansion introduces an explicit time-dependence into the action of any quantum fields present.  Nevertheless, the effective theory idea is a universal one which is not restricted to a specific set of symmetries.  As long as a clear separation of the natural dynamical scales exists within a system, and as long as the theory is being applied to measurements at energies well below the higher of the scales, an effective description should exist even when there is an explicit time-dependence.  Before explaining what is new for such settings, it is important to remember how effective field theories are ordinarily applied to scattering experiments \\cite{Georgi:1994qn,Rothstein:2003mp,Burgess:2007pt}.  A quantum field is introduced for each particle whose mass is below the maximum energy available to an experiment, and all of the symmetries are identified.  These symmetries typically include the set of Poincar\\'e transformations plus whatever other symmetries might be relevant---gauge symmetries, chiral symmetries, permutation symmetries, etc.  The action for the effective theory is then given by the set of all the local operators built from these fields that are invariant under the given symmetries.  This prescription generates an infinite set of operators, though the fact that any given experimental measurement is of finite precision means that in practice only a finite number of them are needed, as long as the energies and momenta of the particles are not too large.  An effective field theory of this sort is never meant to be a fundamental theory, final and absolute, but only one that is appropriate up to a limit.  This limit is usually expressed as a mass scale $M$, and the theory remains predictive for processes whose energies and momenta are below $M$.  But this scale has a deeper meaning too.  Any effective theory always recognizes the possibility that there are further heavier dynamics which cannot be probed directly at low energies.  So above this scale $M$, there should be new fields and interactions, whose masses were too heavy to have had been seen at lower energies.  This higher energy theory is not a final theory either.  Once we have identified all of its fields and symmetries, the game begins anew:  we construct another effective theory which is applicable up to some still higher mass scale where some yet further dynamics remain to be discovered.  If we know the deeper theory or if we have a conjecture for what it might be, we can use it to explain the origin of the low energy effective theory.  Both theories should make the same predictions for the same physical processes.  If we evaluate a process in the higher energy theory, in which external light particles are exchanging virtual heavy particles, the low energy effective theory must yield the same result even though it has no heavy particles to exchange.  This matching of the predictions of the two theories is accomplished by including the appropriate set of interactions in the effective theory.  They have the same external structures in light fields, but unlike the deeper theory they are without any internal structure.  As Feynman graphs, lines representing a heavy field in the higher energy theory are shrunk to a point in the corresponding graphs of the effective theory.  The fact that an action composed of local operators of the relevant light fields can describe any low energy experiment is an important insight underlying effective field theories. It is rooted in the decoupling theorem \\cite{Appelquist:1974tg}, which states that any heavy fields that we have left out of our effective theory will influence the dynamics at low energies only indirectly, by inducing or renormalizing the local operators of the light fields.  The processes that are produced by graphs containing internal heavier particles in the high energy theory are reproduced by an infinite set of local operators in the effective theory.  Seen thus, the coefficients of all of the local operators in the effective theory are completely fixed by the parameters of the higher energy theory.  This relation means that even before we have sufficient energies to probe its dynamics directly, we can infer something about the deeper theory from the sizes and relations amongst the coefficients of the effective theory.  This in brief is the picture for how effective theories are used to analyze particle scattering experiments.  There the relevant physical quantity is the amplitude for one set of initial particles to pass into another set of final particles, which is described by the scattering matrix, or $S$-matrix.  The detailed time-evolution is not so important as the asymptotic inputs and results, and the action is invariant under time-translations as a part of a larger Poincar\\'e invariance.  While this is an appropriate symmetry for the times and distances over which particle interactions occur in a scattering experiment, in other quantum systems some of the basic space-time symmetries are explicitly broken.  An important example is the theory of inflation.  In inflation, the field responsible for an accelerated expansion of the universe also undergoes quantum fluctuations.  These fluctuations are an essential part of the picture.  They are meant to provide the initial spatial inhomogeneities in the universe that eventually grow into all of the structures that we see.  But unlike particle theories, the action for this quantum field contains an explicit dependence on time inherited from the expanding background.  The physical observables that we are calculating are different too:  we should like to know how the expectation values of products of the fluctuations evolve from some initial state, rather than how initial and final states overlap as in a scattering experiment.  An explicit time-dependence is a generic property of any quantum system that is not in an equilibrium state.  In a condensed matter system, for example, we could raise or lower the temperature, subject it to external fields, and then watch how it responds or relaxes.  Even in particle physics, in the theory of baryogenesis for example, sometimes a stage of out of equilibrium evolution is an essential ingredient.  This article explores how the effective theory idea can be applied to time-dependent systems.  More precisely, we shall try to understand how to treat systems that start in a simple nonequilibrium state and evolve in time.  Since we are investigating time-dependent settings, typically with initial times and finite evolutions, we shall be following the time-evolution of the expectation values of observables rather than calculating scattering processes using the $S$-matrix.  In some settings without Poincar\\'e invariance, as in the case of quantum fields in de Sitter space for instance, there is not even a well defined $S$-matrix \\cite{Witten:2001kn}, clearly indicating the shortcomings of this observable in time-dependent situations.  The appropriate formalism for following the evolution in a time-dependent setting was first developed by Schwinger and Keldysh \\cite{Schwinger:1960qe,Keldysh:1964ud,Bakshi:1962dv,Bakshi:1963bn}.  It allows us to choose arbitrary initial states and it automatically maintains a consistently causal evolution of any observable quantity.  While these methods have been often applied to nonequilibrium quantum systems (see for example \\cite{2004cond.mat.12296K}), and have more recently begun to be more widely used for inflationary models, they are perhaps less familiar to many effective field theorists, so in the next section we review the Schwinger-Keldysh formalism.  Most of this article studies a particular example of a time-dependent theory consisting of a  single light field and a single heavy field playing the roles of a heavy and light sector.  These fields interact with each other and, more importantly, their action includes an explicit dependence on the initial time.  An effective theory, even a time-dependent one, must reproduce the same physical predictions at low energies as the higher energy theory.  So as a minimal prescription we demand that at energies well below the scales associated with the heavy field, $$ \\left\\{ { \\hbox{expectation values of operators} \\atop\\hbox{in the full theory}} \\right\\} = \\left\\{ { \\hbox{expectation values of operators} \\atop\\hbox{in the effective theory}} \\right\\} , $$ where the operators are made up only of the light field.  This matching condition can be taken to define what we mean by an effective theory, given a particular higher energy theory.  One reason for examining a system with a particular time-dependence, rather than considering the most general case from the start, is that once we have broken the time-translation invariance, we could have, in principle, arbitrarily time-dependent structures in the action of the effective theory.  By restricting to a specific case, we can learn whether certain forms of time-dependence translate into generic structures in the low energy theory---or what an effective description even means in this context.  For the simple system that we have investigated, we have found that the effective description contains in part all of the local Poincar\\'e invariant operators consistent with a few additional symmetries, and in part some nonlocal transient operators too, related to the initial state that we chose.  At leading nontrivial order in the strength of the coupling between the two fields, these two types of operators are sufficient to describe the effective theory.  Moreover, the transient operators have a universal time-dependence in their leading rate of decay.  This means that with sufficient time, these operators decay to a level below the threshold that would be detectable for a particular experimental precision, and we are left at this order with only the same local action that would have been found in an $S$-matrix derivation of the effective theory.  But at higher orders, we have found that other effects also appear in the expectation values evaluated in the full theory, which are more difficult to express as having arisen from operators in an action for the effective theory.  In the next section we summarize the tools needed for our analysis:  the Schwinger-Keldysh formalism and the tadpole method used to construct the equation of motion of the expectation value of a field.  We then apply these to our model and use the effective equation of motion for the expectation value of the light field to infer the types of terms needed in the effective action to reproduce this equation.  We conclude with a discussion of the reasons for the appearance of the nonlocal terms and point out further applications of our results. ",
        "conclusions": "\\label{sec:conc}  In this article we have explored the construction of effective field theories for time-dependent, nonequilibrium systems. \u00caOur basic tool for analyzing the possible structures that could appear in an effective theory was to take a top-down approach, starting with a simple theory with two scalar fields, one heavy and one light, and to demand that the effective theory composed only of the light field should reproduce the observables in a late time and low energy limit. \u00ca  In particular, we have calculated the equation of motion for the expectation value of the light field which includes the leading virtual effects of the heavy field. \u00caWe started the evolution at a finite initial time and chose the initial state to be the vacuum of the free theory. \u00caThe resulting equation of motion includes the usual local terms that would have been expected from a local, Poincar\\'e-invariant action for the effective field theory, but we found that further terms are present which cannot arise from a local effective field theory.  Among them are some with an explicit time dependence, while others contain time integrals of powers of the light field. \u00caIn the limit where the initial time is taken to the infinite past, and the coupling is turned on adiabatically, the nonlocal terms all vanish and a local effective theory is recovered.  If it appears that the action that we found is not what one would have expected of a theory when the heavy and light sectors are decoupled, which would have produced only the set of local operators, it is because our example essentially contains excitations of the heavier field.  The initial state, which was chosen to be the vacuum for the {\\it free\\/} Lagrangian of both fields, is not an eigenstate of the interacting theory.  So from the very start, we are not in an equilibrium state.  Once we have opened up the system thus, the evolution no longer needs to be unitary.  In our example, the nonlocal terms correspond to the annihilations of the excitations of the heavy field contained in our initial state.  Another way of seeing the origin of this nonlocality is to extend the evolution back to $-\\infty$, and to reproduce the original setting by introducing the interaction as ${1\\over 2}g\\Theta(t-t_0)\\chi^2\\Phi^2$, which explicitly breaks the time-translation invariance.  We can do so since our initial states are the free vacuum states (we are here neglecting the self-interactions of the field), and by having the free theory prior to $t_0$ the fields remain in their vacua.  Of course this trick is a heuristic one, since all that we need to know is the action and the initial state:  everything prior to $t_0$ is irrelevant for the subsequent evolution.  But seen as a term that turns on suddenly---and being a step-function it is sudden from the perspective of the heavy field too---we are introducing a finite energy density into the system at $t_0$ which is spread over all momentum modes, including heavy ones.  So the signature of some residual nonlocality as the heavy fields annihilate and the system relaxes is not in itself surprising.  But what we have found is that among some of these nonlocal effects there is a universal time-dependence, a characteristic $-{3\\over 2}$ decay inherited from the structure of the heavy loop.  The existence of these decaying operators means that the rules change slightly.  For a local operator suppressed by $M^{-n}$, if we perform an experiment with a maximal energy of $E$ and a precision of $\\varepsilon$, then our predictions should still take account of all the operators where $(E/M)^n \\gtrsim \\varepsilon$.  But for a class of the nonlocal operators, we must also keep those where $$ {1\\over [M(t-t_0)]^{(3/2)+n'}} \\biggl( {E\\over M} \\biggr)^n \\gtrsim \\varepsilon \\qquad\\quad n,n'=0,1,2,3, \\cdots . $$ So the set of nonlocal operators that we need varies with time.  If we wait long enough---and what we mean by `long enough' depends on the energies and precision of our experiment---then we do not need any of them.  If we make optimal use of our experiment, so that $(t-t_0)^{-1}\\sim E$, then all we have done is to separate out a specific power of $E$ from the usual rule.  However, it is a nonanalytic power, and that would not have been expected of an entirely local effective theory.  The time-dependent decay is slow enough that even when we are below the threshold for exciting the heavier fields directly from the light fields, there is a long enough tail in the annihilation that they are still producing light fields during the regime $1\\ll M(t-t_0)\\ll \\infty$, which might be observable depending on the sensitivity of our experiment.  The relaxation is not `instantaneous' according to the resolution of our detector.  This example is fairly specific, but we expect that there are general lessons which we can learn from it.  Had we started in a more general nonequilibrium state, as long as some heavy modes are excited, we should expect to find similar structures:  nonlocal operators, decaying coefficients.  We must, of course, not include {\\it too large\\/} an energy density in the initial state if we wish to be able to apply an effective treatment.  From the alternative perspective where we turn on the coupling, we could equivalently consider more general time-dependences, ${1\\over 2} g(t)\\chi^2\\Phi^2$; but cases where $g(t)$ has continually occurring features that are appreciable on intervals $\\Delta t\\sim M^{-1}$ would lie outside an effective treatment altogether.  In the opposite regime, we could consider a time-dependent system in which the local effective theory remains always applicable.  For instance, considering a time-dependent coupling again of the form ${1\\over 2} g(t)\\chi^2\\Phi^2$, if we turn on $g(t)$ sufficiently slowly---though not necessarily adiabatically---so that we are only exciting modes below $M$, then we should not expect to see the sorts of nonlocal effects that we found when we excite the heavy modes directly.  Of course, the system is still a time-evolving, nonequilibrium one, where energy can be exchanged amongst different momentum modes of the light field.  We saw an example of this sort of time-dependence in the equation of motion for the light field back when we introduced the tadpole method.  But as long as none of the heavy fields are excited directly, their existence should only enter the theory through local operators of the light field.  These conjectures about the form of the effective theory in the presence of more general time-dependent structures still need to be checked rigorously.  And even for the comparatively simple example that we have analyzed here, a few open questions remain which also must be addressed, such as the fate or role of the terms which appear at order $g^3$ and above that contain integrals of powers of the light field \\cite{later}.  One route for analyzing the structure of the effective theory more directly, at the level of the path integral rather than through the tadpole method used here, would be to express the virtual effects of the heavy field through an {\\it influence functional\\/} \\cite{Feynman:1963fq}.  By appropriately expanding this influence functional in the limit where the momenta entering it are small compared to $M$, we could see what structures appear in the path integral.  From this we could then try to infer the appropriate generating function for the effective theory.  The appearance of time-dependent structures in the high energy action, such as the coupling here that `turns on' suddenly or some more general coupling with an ongoing time-dependence, represent vertices that do not conserve energy.  Analyzing these vertices in the time-frequency domain, we can see exactly which frequencies of which fields are being excited by a particular time-dependent operator.  Ultimately we should like to apply what we have learned to systems where there is a natural time-dependence, especially in instances where the higher energy theory might be unknown.  In inflation, there is a rather dramatic time-dependence produced by the expansion of the space-time.  The existence of the oscillations in the nonlocal, decaying terms might then be able to produce interesting features in the spectrum of primordial fluctuations, for example.  Some effort has already been made to apply the ideas of effective field theories to inflationary backgrounds \\cite{Weinberg:2005vy,Weinberg:2008hq,Cheung:2007st,Senatore:2009cf}.  And some works too have specifically examined the influence of a heavy field on the correlation functions of the inflaton \\cite{Achucarro:2012sm,Achucarro:2010da,Shiu:2011qw,Avgoustidis:2012yc}, using classical field theory.  The picture that is emerging is a very intriguing one.  It is clear that there is still much to be learned from effective theories in time-dependent settings."
    },
    "1208/1208.2910_arXiv.txt": {
        "abstract": "{ The frequent presence of weak magnetic fields on the surface of spotted late-B stars with HgMn peculiarity in binary systems has been controversial during the two last decades. Recent studies of magnetic fields in these stars using the least-squares deconvolution (LSD) technique have failed to detect magnetic fields, indicating an upper limit on the longitudinal field between 8 and 15\\,G. In these LSD studies, assumptions were made that all spectral lines are identical in shape and can be described by a scaled mean profile. } { We re-analyse the available spectropolarimetric material by applying the moment technique on spectral lines of inhomogeneously distributed elements separately. Furthermore, we present new determinations of the mean longitudinal magnetic field for the HgMn star HD\\,65949 and the hotter analog of HgMn stars, the PGa star HD\\,19400, using FORS\\,2 installed at the VLT. We also give new measurements of the eclipsing system AR\\,Aur with a primary star of HgMn peculiarity which were obtained with the SOFIN spectropolarimeter installed at the Nordic Optical Telescope. } { We downloaded from the European Southern Observatory (ESO) archive the publically available HARPS spectra for eight HgMn stars and one normal and one superficially normal B-type star obtained in 2010. Out of this sample, three HgMn stars belong to spectroscopic double-lined systems. The application of the moment technique to the HARPS and SOFIN spectra allowed us to study the presence of the longitudinal magnetic field, the crossover effect, and quadratic magnetic fields. Results for the HgMn star HD\\,65949 and the PGa star HD\\,19400 are based on a linear regression analysis of low-resolution spectra obtained with FORS\\,2 in spectropolarimetric mode. } { Our measurements of the magnetic field with the moment technique using spectral lines of several elements separately reveal the presence of a weak longitudinal magnetic field, a quadratic magnetic field, and the crossover effect on the surface of several HgMn stars as well as normal and superficially normal B-type stars. Furthermore, our analysis suggests the existence of intriguing correlations between the strength of the magnetic field, abundance anomalies, and binary properties. The results are discussed in the context of possible mechanisms responsible for the development of the element patches and complex magnetic fields on the surface of late B-type stars. } {} ",
        "introduction": "\\label{sect:intro}  The origin of abundance anomalies observed in late B-type stars with HgMn peculiarity is still poorly understood. Over the last few years, we have performed extensive spectroscopic studies of both single late B-type stars and of spectroscopic binaries (SB) with late B-type primaries  (spectral types B7--B9) with the goal of understanding why the vast majority of these stars exhibit in their atmospheres certain chemical abundance anomalies, i.e.\\ large excesses of P, Mn, Ga, Br, Sr, Y, Zr, Rh, Pd, Xe, Pr, Yb, W, Re, Os, Pt, Au, and Hg, and underabundances of He, Al, Zn, Ni, and Co (e.g. Castelli \\& Hubrig \\cite{Castelli2004a}). Strong isotopic anomalies were detected for the chemical elements Ca, Pt, and Hg, with patterns changing from one star to the next (Hubrig et al.\\ \\cite{Hubrig1999a}; Castelli \\& Hubrig \\cite{Castelli2004b}; Cowley et al.\\ \\cite{Cowley2010}). Observationally, these stars are characterised by low rotational velocities ($\\langle v\\,\\sin i\\rangle \\leq 29$~km\\,s$^{-1}$, Abt et al.\\ \\cite{Abt1972}). The number of these chemically peculiar stars, usually called HgMn stars, decreases with increasing rotational velocity (Wolff \\&\\ Wolff \\cite{Wolff1974}).  More than two thirds of the HgMn stars are known to belong to spectroscopic binaries (Hubrig \\& Mathys \\cite{Hubrig1995}), with a preference of orbital periods ranging from 3 to 20 days. It is striking that the inspection of SB systems with a late B-type primary in the 9$^{\\rm th}$ Catalogue of Spectroscopic Binary Orbits (Pourbaix et al.\\ \\cite{Pourbaix2009}) indicates a strong correlation between the HgMn peculiarity and membership in a binary system among bright, well-studied SB systems with late B-type, slowly rotating \\mbox{($v\\,\\sin i<\\,70~{\\rm km\\,s}^{-1}$)} primaries with an apparent magnitude of up to $V\\approx7$  and orbital periods between 3 and 20\\,d. With the exception of HR\\,7241, all 21~systems have a primary with a HgMn peculiarity. Based on this fact, it is very likely that the majority of slowly rotating late B-type stars formed in binary systems with certain orbital parameters become HgMn stars. Consequently, careful studies of these peculiar stars are important for the general understanding of B-type star formation in binary systems. According to Abt \\& Snowden (\\cite{Abt1973}), close-binary formation is inhibited when strong magnetic fields are present. They also suggest that close binaries dissipate the magnetic field.  The aspect of inhomogeneous distribution of some chemical elements over the surface of HgMn stars was first discussed by Hubrig \\& Mathys (\\cite{Hubrig1995}). From a survey of HgMn stars in close SBs, it was suggested that some chemical elements might be inhomogeneously distributed on the surface, with, in particular, preferential concentration of Hg along the equator. The first definitively identified spectrum variability was reported for the binary HgMn star $\\alpha$~And by Wahlgren et al.\\ (\\cite{Wahlgren2001}) and Adelman et al.\\ (\\cite{Adelman2002}), who showed that the spectral variations of the \\ion{Hg}{ii} line at $\\lambda$3984 discovered in high-dispersion spectra are not due to the orbital motion of the companion, but produced by the combination of the 2.8\\,d period of rotation of the primary and a non-uniform surface distribution of mercury, which is concentrated in the equatorial region. Their results are in good correspondence with those of Hubrig \\& Mathys (\\cite{Hubrig1995}). The variability of the \\ion{Hg}{ii} line at $\\lambda$3984 was interpreted with a Doppler imaging code, revealing high-contrast mercury spots located along the rotational equator. Using a Doppler imaging reconstruction of spectroscopic time series obtained over seven consecutive years, Kochukhov et al.\\ (\\cite{Kochukhov2007}) suggested the presence of a secular evolution of the mercury distribution. On the other hand, the work of other authors has proved that not only mercury abundance appears distributed in patches over the stellar surface. Almost all other elements, most typical Ti, Cr, Fe, Mn, Sr, Y, and Pt, are concentrated in spots of diverse size, and different elements exhibit  different abundance distributions across the stellar surface (e.g. Hubrig et al.\\ \\cite{Hubrig2006a}; Briquet et al.\\ \\cite{Briquet2010}; Makaganiuk et al.\\ \\cite{Makaganiuk2011a}). Moreover, an evolution of the element abundance spots at different time scales was discovered in two additional HgMn stars.  Briquet et al.\\ (\\cite{Briquet2010}) reported the presence of dynamical spot evolution over a couple of weeks for the SB1 system HD\\,11753, while Hubrig et al.\\ (\\cite{Hubrig2010}) detected a secular element evolution in the double-lined eclipsing binary AR\\,Aur. Importantly, recent results (e.g. Nu{\\~n}ez et al.\\ \\cite{Nunez2011}; Hubrig et al.\\ \\cite{Hubrig2011}) show that line-profile variability of various elements caused by a non-uniform abundance distribution is a general characteristic of HgMn stars, rather than an exception.  Typically, inhomogeneous chemical abundance distributions are observed only on the surface of magnetic chemically peculiar stars with large-scale organised magnetic fields (Ap and Bp stars). In these stars, the abundance distribution of certain elements is non-uniform and non-symmetric with respect to the rotation axis. Numerous studies of Ap and Bp stars have revealed a kind of symmetry between the topology of the magnetic field and the element distribution. Thus, the structure of the magnetic field can be studied by measurements of the magnetic field, using spectral lines of each element separately. Although strong large-scale magnetic fields have not generally been found in HgMn stars, it has never been ruled out that these stars might have tangled magnetic fields of the order of a few thousand Gauss with only very weak net longitudinal components. Bychkov et al.\\ (\\cite{Bychkov2009}) compiled all published longitudinal magnetic field measurements of different groups of chemically peculiar stars and concluded that the group of HgMn stars possesses the second weakest fields after the group of Am stars.  In the last few years, a number of  attempts to detect mean longitudinal magnetic fields in HgMn stars have been made by several authors using the line addition technique, the least-squares deconvolution (LSD), most recently by Makaganiuk et al.\\ (e.g. \\cite{Makaganiuk2011a}, \\cite{Makaganiuk2011b}, \\cite{mak2011c}). A high level of precision, from a few to tens of Gauss, is achieved through application of the LSD technique (Donati et al.\\ \\cite{Donati1997}), which combines hundreds of spectral lines of various elements. This technique assumes that all spectral lines are identical in shape and can be described by a scaled mean profile. However, the lines of different elements with different abundance distributions across the stellar surface sample the magnetic field in different manners. Combining them as is done with the LSD technique may lead to the dilution of the magnetic signal or even to its (partial) cancellation, if enhancements of different elements occur in regions of opposite magnetic polarities. Using this technique, Makaganiuk et al.\\ (\\cite{Makaganiuk2011b}) analysed HARPS spectra for a sample of HgMn and normal B-type stars and reported no detection at a 3$\\sigma$ level in any of the studied targets.  Strangely enough, although the authors were aware of the inhomogeneous distribution of the elements on the surface of HgMn stars, no analysis has been done on the lines of individual elements separately. Only in the most recent Makaganiuk et al.\\ (\\cite{mak2012}) study did the authors decide to constrain the longitudinal magnetic field of HD\\,11753 by computing mean profiles for Y, Ti, and Cr, which show different spot distributions on the stellar surface. Their results indicate that the upper limit for the strength of the magnetic fields can reach about 20--30\\,G, with a typical error bar from 8\\,G to 15\\,G.  \\begin{table} \\centering \\caption{ List of the studied targets. } \\label{tab:targetlist} \\centering \\begin{tabular}{lcrcc} \\hline \\hline \\multicolumn{1}{c}{Object \\rule{0pt}{2.6ex}} & \\multicolumn{1}{c}{Other} & \\multicolumn{1}{c}{V} & \\multicolumn{1}{c}{Spectral} & \\multicolumn{1}{c}{Instr.} \\\\ \\multicolumn{1}{c}{Name} & \\multicolumn{1}{c}{Identifier} & & \\multicolumn{1}{c}{Type} & \\\\ \\hline HD\\,11753     & $\\phi$\\,Phe & 5.11 & B9V   & HARPS \\\\ HD\\,19400     & $\\theta$\\,Hyi & 5.50 & B3V   & FORS\\,2 \\\\ HD\\,27376     & 41\\,Eri     & 3.55 & B9V, SB2   & HARPS \\\\ HD\\,32964    & 66\\,Eri       & 5.10  &  B9V, SB2    & HARPS \\\\ HD\\,33904   & $\\mu$\\,Lep    & 3.28  & B9IV     & HARPS \\\\ HD\\,34364   & AR\\,Aur        & 6.14  & B9V, SB2    & SOFIN \\\\ HD\\,53244   & $\\gamma$\\,CMa   & 4.10  & B8II    & HARPS \\\\ HD\\,65949   & CPD$-$60\\,966    & 8.37  & B8/B9    & FORS\\,2 \\\\ HD\\,78316   & $\\kappa$\\,Cnc   & 5.24  & B8III, SB2    & HARPS \\\\ HD\\,101189   &HR\\,4487        & 5.14  & B9IV    & HARPS \\\\ HD\\,221507   & $\\beta$\\,Scl   & 4.37  & B9.5IV    & HARPS \\\\ \\hline \\multicolumn{5}{c}{(Superficially) normal B-type stars}\\\\ \\hline HD\\,179761  &V1288\\,Aql & 5.14 & B8II-III     & HARPS \\\\ HD\\,209459  & 21\\,Peg    & 5.82 & B9.5V      & HARPS \\\\ \\hline \\end{tabular} \\tablefoot{ Spectral types and visual magnitudes are taken from SIMBAD. } \\end{table}  To test if the use of the LSD technique may indeed account for the non-detection of magnetic fields in HgMn stars, we decided to re-analyse the HARPS spectropolarimetric material that recently became publically available in the ESO archive. In this work, we use a completely different approach for the measurements of magnetic fields, namely, the moment technique developed by Mathys (e.g. \\cite{Mathys1991}, \\cite{Mathys1995a}, \\cite{Mathys1995b}). This technique allows us not only to determine the mean longitudinal magnetic field, but also to prove the presence of the crossover effect and quadratic magnetic fields. This information cannot be obtained from the LSD technique, as it assumes  that all spectral lines are identical in shape and can be described by a scaled mean profile. Furthermore, we present five new spectropolarimetric observations of the eclipsing system AR\\,Aur with a primary star of HgMn peculiarity, obtained with the echelle spectrograph SOFIN at the Nordic Optical Telescope at the end of 2010. The variability and spot distribution of this star has been intensively studied by our group in the few last years (Hubrig et al.\\ \\cite{Hubrig2006a}, \\cite{Hubrig2011}). In addition, a new set of polarimetric spectra was obtained for  the HgMn star HD\\,65949 and the hotter analog of HgMn stars, the PGa star HD\\,19400, using FORS\\,2 installed at the VLT. All targets discussed in this work are presented in Table~\\ref{tab:targetlist} together with their visual magnitudes and spectral types.  In Sect.~\\ref{sect:obs}, we describe the observations, data reduction, and methods of our magnetic field measurements. The stellar characteristics and the  magnetic field measurements for the individual targets are reviewed in Sect.~\\ref{sect:indivi}. Finally, in Sect.~\\ref{sect:disc}, we discuss our results in the context of possible mechanisms at play, which can be considered responsible for the development of the inhomogeneous abundance distribution and the presence of weak magnetic fields. ",
        "conclusions": "\\label{sect:disc}  The region of the main sequence centred on A and B stars, also referred to as the ``tepid stars'', represents an ideal laboratory to study a wide variety of physical processes that are at work to a greater or lesser extent in most stellar types. These processes include radiation driven diffusion, differential gravitational settling, grain accretion, magnetic fields and non-radial pulsations. Understanding them is becoming increasingly important for the refinement of stellar evolution models and for the improved treatment of the stellar contribution in studies of galactic evolution. Presently, the radiative diffusion hypothesis developed by Michaud (\\cite{Michaud1970}) is the most frequently accepted theory for producing the observed anomalies in HgMn stars. According to Michaud et al.\\ (\\cite{Michaud1974}), the effect of diffusion would cause a cloud of mercury to form high up in the HgMn star atmospheres. However, the impact of the radiative diffusion process has not yet been studied for any element considered in our work.  Based mostly on ESO/HARPS archive data, our analysis suggests the existence of intriguing correlations between magnetic field, abundance anomalies, and binary properties. In the SB2 systems with the synchronously rotating components, 41\\,Eri and AR\\,Aur, the stellar surfaces facing the companion star usually display low-abundance element spots and negative magnetic field polarity. The surface of the opposite hemisphere, as a rule, is covered by high-abundance element spots and the magnetic field is positive at the rotation phases of the best-spot visibility. Also, Measurement results for the SB1 system HD\\,11753 indicate that element underabundance (respectively overabundance) is observed where the polarity of the magnetic field is negative (respectively positive). In the case of the very young pseudosynchronously rotating system 66\\,Eri, additional future observations are urgently needed to determine the period of the spectral variability and to investigate the presence of element-spot evolution. It is quite possible that because the system is not yet synchronised, a dynamical spot evolution is more pronounced compared to synchronised systems.  Although only rather weak longitudinal magnetic fields are detected on the surface of our targets, the numerous 3$\\sigma$ detections of mean quadratic magnetic fields strongly suggest that magnetic fields are present in their atmospheres. We note again the advantage of carrying out quadratic magnetic field measurements as they provide measurements of the magnetic field strength that are somewhat insensitive to the magnetic field structure, and especially suited for magnetic fields with complex structure.  The interactions of differential rotation, magnetic fields, and meridional flows in a tidal potential may lead to processes that can explain complex surface patterns. Possible instabilities should be studied by non-linear, three-dimensional simulations. The resulting fluctuations in velocity and magnetic field may not only explain the relatively small scales of surface magnetic fields, but also provide enhanced transport of angular momentum that explain the high rates of synchronisation. Synchronisation in an orbit with about 5\\,d period, an initial rotation period of 1\\,d , and stellar parameters similar to 66~Eri is roughly $10^4$ times the dissipation time scale in the star (Zahn \\cite{Zahn2008}). In a non-turbulent, non-magnetic star where only the microscopic gas viscosity can couple the surface with the interior, it would be impossible to reach synchronisation within the age of the universe. There are also gravity waves that lead to synchronisation, but for stars of 2--3~solar masses, this effect still takes roughly $10^9$~years. Large-scale magnetic fields or turbulence can provide enough coupling to reach (pseudo-)synchronisation in about a million years.  \\begin{figure} \\centering \\includegraphics[angle=270,width=0.45\\textwidth]{e-t.eps} \\caption{ Circularisation times for 41\\,Eri (dotted lines), 66\\,Eri (solid lines),  and AR\\,Aur (dashed lines), assuming that they formed with eccentric orbits with $e_0$ = 0.2 and 0.5. } \\label{fig:et} \\end{figure}  As an illustration, we calculated the circularisation times for the systems 41\\,Eri, 66\\,Eri, and AR\\,Aur, assuming that they formed with eccentric orbits. With the present parameters of the stars, the circularisation time scales, $ (-d\\ln e/dt)^{-1}$ could be as high as $10^{10}$ yr. As they evolve within the main sequence, circularisation becomes more efficient, but even so they will arrive at the terminal age main sequence (TAMS) with a significant fraction of their original eccentricity: $e/e_0$=0.15--0.70, depending on the original eccentricity $e_0$. Figure~\\ref{fig:et} shows the evolution of $e$ for the three systems, starting from $e_0$ = 0.2 and 0.5. The solid line corresponds to 66\\,Eri, the dotted line to 41\\,Eri, and the dashed line to AR\\,Aur. The  vertical end of the curves to the right mark the TAMS. The calculations were performed using the stellar models calculated by Claret (\\cite{Claret2004}), which include the internal structure constants required to calculate the orbital evolution. Considering that all three systems are young, with ages well below $10^8$ yr, it would be expected that their present eccentricities woould not differ significantly from the original ones. The observed very low or null eccentricities cannot be explained as a consequence of standard tidal evolution.  The magneto-rotational instability is a candidate that can provide turbulence for a wide range of magnetic field strengths and field topologies, but only if the internal angular velocity of the star decreases with the distance from the rotation axis. Since synchronisation in the systems considered here is in fact a braking process, the star may exhibit a slower rotation near its surface than deeply inside. Numerical simulations should elucidate the expected enhancement of the viscosity in the star and will also deliver surface topologies of the fields resulting from this or other instabilities that can be compared to the observed patterns.  Tidal forces strongly depend on the distance between the components. For the three SB2 systems with synchronised and pseudosynchronised  components, 41\\,Eri, 66\\,Eri, and AR\\,Aur, we estimated the fractional radii $r_j=R_j/a$. Assuming for 41\\,Eri  \\begin{eqnarray} (v\\,\\sin\\,i)_A\t&=& 12.23\\pm{}0.06\\,{\\rm km/s}, \\nonumber \\\\ (v\\,\\sin\\,i)_B &=& 11.78\\pm{}0.07\\,{\\rm km/s}, \\nonumber \\\\ K_A+K_B &=& 127.38 \\pm 0.23\\,{\\rm km/s}, \\nonumber \\\\ e&=&0,\\nonumber \\end{eqnarray} it follows that $R_A/r = 0.096\\pm{}0.001$ and $R_B/r = 0.093\\pm 0.001$  If we assume for 66\\,Eri pseudosynchronisation ($P_{\\rm rot} = 5.3015\\,d$), we obtain $R_A/r = 0.084\\pm 0.001$ and  $R_B/r = 0.080\\pm 0.001$.  For AR\\,Aur we use \\begin{eqnarray} (v\\,\\sin\\,i)_A &=& 22.87\\pm 0.08\\,{\\rm km/s}, \\nonumber \\\\ (v\\,\\sin\\,i)_B &=& 22.35\\pm{}0.07\\,{\\rm km/s}, \\nonumber \\\\ K_A+K_B &=& 225.28 \\pm 0.18\\,{\\rm km/s}, \\nonumber \\\\ e &=& 0 \\nonumber \\end{eqnarray} and obtain  $R_A/a = 0.102\\pm 0.001$ and $R_B/a = 0.099\\pm 0.001$.\\\\  Since the fractional radii in these systems are quite large, the components in the close binary systems are subject to tidal forces acting differentially throughout their bodies and causing various effects in the system. The tidal effects include synchronisation of the rotational periods with the orbital period, circularisation of the orbit, and alignment of the rotation axes with the orbital axis. All these changes in the dynamical properties of the stars happen on different time scales. The spin-down may be a surface effect, but the star then possesses an internal differential rotation that may lead to hydrodynamic or magnetohydrodynamic instabilities. Even a moderately turbulent state generated by an instability will lead to much stronger coupling with the stellar interior and redistribute angular momentum.  \\begin{figure*} \\centering \\includegraphics[angle=270,totalheight=0.55\\textwidth]{hr4.eps} \\caption{H-R diagram for the HgMn stars in our sample. Continuous lines present the ZAMS and TAMS. Dotted lines present isochrones for log(t) from 7.00 to 8.70. } \\label{fig:HR} \\end{figure*}  To establish the evolutionary status of the HgMn stars in our sample, we studied their distribution in the H-R diagram. To this aim, the distances were calculated mainly using Hipparcos parallaxes (van Leeuwen \\cite{hipp}). For stars with high galactic latitudes, we estimated the interstellar absorption using Schlegel et al.\\ (\\cite{Schlegel1998}) maps and the distance by applying the same procedure as Bilir et al.\\ (\\cite{Bilir2008}). Since this approach is not valid for targets located very close to the galactic plane, we estimated the absorption in a different way for the low-latitude stars HD\\,101189 and AR\\,Aur.. HD\\,101189 is located in the field of the open cluster NGC\\,3766, but it is a foreground star. This fact allowed us to estimate the absorption to the star using cluster reddening and the distance ratio. For the eclipsing binary AR\\,Aur, we adopted the absolute stellar parameters derived by Nordstrom \\& Johansen (\\cite{nord1994}) from light and radial velocity curves. The star  HD\\,65949, which is the faintest and most distant star, belongs to the open cluster NGC\\,2516. We used, therefore, the cluster distance and reddening published by Terndrup et al.\\ (\\cite{phot2516}) to obtain absolute magnitude and intrinsic color. In the case of double-lined spectroscopic binaries, the position of the individual components in the H-R diagram was calculated using the spectroscopic mass ratio, from which the magnitude difference between the components was estimated by interpolating in the Geneva stellar models (Schaller et al.\\ \\cite{Schaller1992}). Figure\\,\\ref{fig:HR} shows the position of the 11 HgMn stars in the color-magnitude diagram, together with isochrones taken from Schaller et al.\\ (\\cite{Schaller1992}). Six stars in our sample, 41\\,Eri\\,A, 41\\,Eri\\,B, 66\\,Eri\\,A, AR\\,Aur\\,A, HD\\,65949, and HD\\,221507, are located very close to the ZAMS.  From the knowledge of orbital parameters, fundamental parameters of components, and fractional radii, it is possible to estimate the reflection effect, gravitational distortion, and darkening for a typical HgMn star in a SB2 system. One of the consequences of the proximity of the stars is the departure from sphericity of the stellar surface. The star acquires a somewhat larger dimension along the line passing through the two star centres, while it has a smaller radius in the direction that is perpendicular to the former line and also lies in the orbital plane. The relative differences between both radii are 0.09\\% and 0.08\\% for the two components of the system 41\\,Eri, 0.05\\% and 0.06\\% for the system 66\\,Eri, and 0.09\\% and 0.10\\% for the system AR\\,Aur. Observational effects of such stellar shapes are, for example, the variation of the $v\\sin i$-value, which in the three systems studied is below the observational errors, and light variations due to the variation of the stellar disk size, which is expected to be of the order of 2--3\\,mmag. An effect that could affect physical conditions in the stellar atmosphere is the reflection effect. Assuming an albedo equal to 1 (the usual assumption for stars with radiative envelopes) and the temperatures and geometrical configuration of each system, we estimate that the temperature excess in the surface region facing the companion is between 20 and 40 K, which represents 0.18--0.38\\% of the star temperature. The question to solve is whether such small temperature differences can be considered a reasonable explanation of the origin of the low- and high-abundance spots observed either on the surface facing the companion or on the surface hidden to the companion. Although it is currently not easy to give an answer without a specific model of time-dependent diffusion processes, it seems very unlikely that these processes are the main cause of the observed abundance pattern. We note that the early circularisation and a very specific configuration of chemical spots in binary systems support the idea of an important impact of magnetic fields on the physical processes taking place in these stars. Since the orbital parameters of close binaries can change as a consequence of the action of tidal forces or other mechanisms, the influence of the presence of a stellar companion on the physical processes playing a role in the formation of chemical anomalies is expected to vary with time. The atmospheric physical conditions for the development of chemical inhomogeneities might appear favourable at certain stages of the stellar life and unfavourable at other stages.   Further studies using high S/N circularly polarised spectra well-distributed throughout the rotation period of a representative sample of  HgMn spectroscopic binaries sampling a wide range  of inclinations of the rotation axis to the line of sight are urgently needed to map their magnetic field and elemental abundances with a Zeeman-Doppler imaging code. The sampling of a wide range  of inclination of the rotation axis to the line of sight in HgMn systems is especially important because the latitudinal information on the surface abundance distribution of different elements in HgMn stars is still rather poor. In this way, it will be possible to probe correlations between the binary properties, the magnetic field structure, and the abundance inhomogeneities. At the same time, we will be able to gain insight into the mechanisms responsible for the chemical inhomogeneities and their dynamical evolution on the surface of HgMn stars.  {"
    },
    "1208/1208.0638_arXiv.txt": {
        "abstract": "The \\bicep2 and \\keck\\ experiments are designed to measure the polarization of the cosmic microwave background (CMB) on angular scales of 2-4 degrees ($\\ell=50$--100).  This is the region in which the $B$-mode signal, a signature prediction of cosmic inflation, is expected to peak. \\bicep2 was deployed to the South Pole at the end of 2009 and is in the middle of its third year of observing with 500 polarization-sensitive detectors at $150~\\mathrm{GHz}$.  The \\keck\\ was deployed to the South Pole at the end of 2010, initially with three receivers---each similar to \\bicep2.  An additional two receivers have been added during the 2011-12 summer.  We give an overview of the two experiments, report on substantial gains in the sensitivity of the two experiments after post-deployment optimization, and show preliminary maps of CMB polarization from \\bicep2. ",
        "introduction": "\\label{sec:intro}  % Over the past 15 years, the $\\Lambda$CDM theory has emerged as the ``standard model'' of cosmology. It describes a universe in which cold dark matter drives the formation of structure, and dark energy (or a cosmological constant) causes accelerating expansion of space. This concordance model explains the observed universe well, but leaves several important questions unanswered, such as the nature of the dark energy and dark matter. It accommodates, but does not predict, the observed flatness, homogeneity, and isotropy of the universe. These three properties can be explained by the addition to the model of cosmic inflation.  If the universe went through an inflationary phase in the first instant after the Big Bang, this rapid, exponential expansion from a single microscopic volume would have smoothed out spatial curvature.  The observed homogeneity and isotropy arise from the fact that the entire observable universe would have been in causal contact before inflation.  Quantum fluctuations in the inflaton field give rise to scale-invariant density fluctuations that would seed structure formation and the primary temperature anisotropy of the cosmic microwave background (CMB). Current CMB observations are in good agreement with inflationary predictions. \\cite{boomerang00,wmapinfl03}  In addition to these \\emph{scalar} density perturbations, inflation also generates primordial gravitational waves, i.e. \\emph{tensor} metric perturbations.  Both types of perturbation give rise to temperature variations at the time when CMB photons last scattered, but with different spatial structure; they accordingly lead to different signatures in the CMB.  They can be distinguished by decomposing the patterns of CMB polarization into two classes: the $E$-modes, which have even parity and arise from scalar and tensor perturbations; and the $B$-modes, which have odd parity and arise from tensor perturbations only. The detection of $B$-modes would confirm a key prediction of inflation, and a measurement of (or upper limit on) the amplitude of $B$-modes would provide information about the energy scale of inflation, which determines the tensor-to-scalar ratio $r$.\\cite{primer97}  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=8cm]{spud5.eps} \\end{tabular} \\end{center} \\caption[example] { \\label{fig:spud5} \\keck\\ telescope in February 2012, with five receivers installed in the mount originally built for \\dasi\\ at the South Pole. } \\end{figure}  The \\dasi\\ experiment measured degree-scale $E$-modes in 2002, the first detection of CMB polarization.\\cite{dasi02} The amplitude of $B$-modes is expected to be lower, and their detection remains an open challenge.  The \\bicep2, and \\keck\\ experiments have been designed to meet this challenge. The two experiments adopt elements from the successful \\bicep\\ telescope, which observed from 2006--08.  \\bicep\\ has set the most sensitive limits on $B$-mode polarization of the CMB in the range $\\ell=21-335$, and constrained $r<0.72$ at 95\\% confidence level.\\cite{chiang10,takahashi10} \\bicep2 and the \\keck\\ use a new detector technology, the Caltech-JPL antenna-coupled TES arrays.\\cite{orlando09,orlando10} These detectors are read out using NIST SQUID amplifiers with the University of British Columbia MCE control and readout electronics.\\cite{mce08} The monolithic fabrication and multiplexed readout allow \\bicep2 and the \\keck\\ to field larger numbers of detectors, for improved sensitivity over their predecessors. \\bicep2 was deployed to the South Pole in late 2009\\cite{ogburn10,aikin10,brevik10} and has observed through 2010 and 2011, with a third year of observations now in progress.  With its compact size and low cost, the \\bicep-style refracting telescope can be scaled up by building an array of multiple receivers.  This is the idea behind the \\keck\\, which uses the same detector technology as \\bicep2 and a similar telescope design, with a few additional improvements.\\cite{sheehy10} Instead of consumable liquid helium, the \\keck\\ uses \\cryomech\\ pulse-tube coolers to provide cooling at $4~\\mathrm{K}$ and above. It uses the larger mount originally built for the \\dasi\\ experiment and attached to the Martin A. Pomerantz Observatory (MAPO).  This mount can accommodate up to five \\bicep2-style receivers. The \\keck\\ was deployed and commissioned in December 2010--January 2011, with three cryostats housing three full focal planes of detectors.  It observed with three receivers through 2011, and was upgraded to five receivers for the 2012 season (Fig.~\\ref{fig:spud5}).  Both experiments observe the CMB in the same primary field as \\bicep1, dubbed the \"Southern Hole,\" with secondary observations of a bright region of the Galactic plane.  In this paper we present an overview of the \\bicep2 and \\keck\\ experimental design; a number of improvements to the sensitivity of each experiment that have been made since deployment; and a look at $E$- and $B$-mode maps from two seasons of \\bicep2 observation. Several companion papers in this volume present % the sensitivity of the \\keck\\ (Kernasovskiy \\etal\\cite{kernasovskiy12}), studies of beam shape and differential pointing (Vieregg \\etal\\cite{vieregg12}), and further developments in the detector fabrication for the \\keck\\ and other experiments (O'Brient \\etal\\cite{obrient12}). ",
        "conclusions": "\\bicep2 has completed two years of observation at the South Pole, and is currently in the middle of its third season.  The \\keck\\ has completed its first season, with three receivers, and is now in the middle of its second season, with five receivers.  Both have made substantial improvements to their sensitivity, through optimization of the detector biases and readout rate, and through the addition of new \\keck\\ receivers and replacement of detectors.  \\bicep2 has accumulated a total of 11320 hours on the primary CMB field. We have presented for the first time preliminary \\bicep2 maps of the primary CMB field in polarization, showing the improvement in sensitivity over the original \\bicep\\ experiment."
    },
    "1208/1208.2584_arXiv.txt": {
        "abstract": "We present new measurements of the mean transmitted flux in the \\lya\\ forest over $2 < z < 5$ made using 6065 quasar spectra from the Sloan Digital Sky Survey DR7.  We exploit the general lack of evolution in the mean quasar continuum to avoid the bias introduced by continuum fitting over the \\lya\\ forest at high redshifts, which has been the primary systematic uncertainty in previous measurements of the mean \\lya\\ transmission.  The individual spectra are first combined into twenty-six composites with mean redshifts spanning $2.25 \\le z_{\\rm comp} \\le 5.08$.  The flux ratios of separate composites at the same rest wavelength are then used, without continuum fitting, to infer the mean transmitted flux, $F(z)$, as a fraction of its value at $z \\sim 2$.  Absolute values for $F(z)$ are found by scaling our relative values to measurements made from high-resolution data by \\citet{fg2008a} at $z \\le 2.5$, where continuum uncertainties are minimal.  We find that $F(z)$ evolves smoothly with redshift, with no evidence of a previously reported feature at $z \\simeq 3.2$.  This trend is consistent with a gradual evolution of the ionization and thermal state of the intergalactic medium over $2 < z < 5$.  Our results generally agree with the most careful measurements to date made from high-resolution data, but offer much greater precision and extend to higher redshifts.  This work also improves upon previous efforts using SDSS spectra by significantly reducing the level of systematic error. ",
        "introduction": "The pattern of absorption lines imprinted on the spectra of distant objects by neutral hydrogen in the intergalactic medium (IGM), known as the ``\\lya\\ forest'', is one of the most fundamental probes of cosmic structure.  The characteristics of this absorption reflect the density distribution, ionization state, and temperature of the intergalactic gas.  The gas, in turn, closely traces the underlying distribution of dark matter, albeit with significant deviations arising from hydrodynamics coupled to the radiative and mechanical feedback from galaxies and AGN.  Consequently, the \\lya\\ forest allows us to probe the baryon physics connecting the evolution of large-scale structure to the highly non-linear processes driving galaxy formation.  A key requirement for utilizing the \\lya\\ forest to study either cosmology or baryon physics is to accurately establish the basic parameters describing the flux distribution.  The most basic of these is the mean transmitted flux, $F = \\langle e^{-\\tau} \\rangle$, and its evolution with redshift.  The mean flux, which is often expressed in terms of an effective optical depth, $\\tau_{\\rm eff}(z) =  -\\ln{F(z)}$, directly constrains, for example, the intensity of the metagalactic ionizing background \\citep[e.g.,][]{rauch1997,mcdonald2001a,bolton2005}.  It also influences statistics such as the flux probability distribution function and power spectrum, among others, which are used to constrain quantities ranging from the temperature of the IGM \\citep[e.g.,][]{schaye2000,zaldarriaga2001,lidz2010,becker2011a} to the free-streaming length of dark matter particles \\citep[e.g.,][]{viel2008}.  \\lya\\ forest studies increasingly rely on comparing real data to artificial absorption spectra drawn from simulations, which must be calibrated to the correct mean flux.  With the increasing availability of high-quality quasar and gamma ray burst spectra taken with large telescopes, and the large number of moderate-resolution quasar spectra gathered with surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{york2000} and the Baryon Oscillation Spectroscopic Survey \\citep[BOSS;][]{dawson2012}, a precise measurement of $F(z)$ is therefore critical for accurately extracting the maximum amount of information from the forest.  The main challenge in measuring \\meanfluxz\\ is to overcome cosmic variance while accurately estimating the unabsorbed continuum of the background objects, which have so far typically been quasars.  Studies using high-resolution (generally $R \\sim 40,000$) spectra have the advantage that the peaks of the \\lya\\ forest transmission should approach the continuum \\citep[e.g.,][]{rauch1997,mcdonald2000,schaye2003,songaila2004,kirkman2005,kim2007,becker2007,fg2008a,dallaglio2008,becker2011a}.  Even at moderate redshifts ($z \\sim2-3$), however, voids in the IGM will have a non-negligible optical depth.  The continuum will therefore lie somewhat above the tops of the transmission peaks, a bias that increases towards higher redshifts.  Alternatively, studies using large sets of moderate-resolution spectra have either attempted to simultaneously solve for \\meanfluxz\\ and the quasar continua in a strictly statistical sense \\citep{bernardi2003}, or have fit continua on an individual basis using principle component analysis \\citep[PCA;][]{mcdonald2005,paris2011} or with bias corrections estimated from model spectra \\citep{dallaglio2009}.  At the highest redshifts ($z > 5$), \\meanfluxz\\ measurements have generally relied on extrapolating a power-law continuum from redward of the \\lya\\ emission line \\citep{songaila2004,fan2006b,becker2007}.  The most accurate \\meanfluxz\\ measurement to date has arguably been made by \\citet{fg2008a}.  These authors used a set of 86 high-resolution quasar spectra and made statistical corrections to the continua as a function of redshift based on artificial spectra drawn from hydrodynamic simulations.  Notably, they identified a possible sharp ($\\Delta z \\lesssim 0.5$) dip in the evolution of the mean opacity near $z \\simeq 3.2$.  A similar feature was found by \\citet{bernardi2003} using a very different data set and measurement technique.  The optical depth of the optically thin IGM to \\lya\\ will scale with the \\hi\\ neutral fraction, which in turn depends on the gas temperature, $T$, and the \\hi\\ ionization rate, $\\Gamma$, as $f_{\\mhi} \\propto T^{-0.7}\\Gamma^{-1}$, where the temperature dependence is for case A recombination.  \\bernardi\\ noted that a  decrease in \\taueffz\\ at $z \\simeq 3.2$ could indicate a temporary increase in the IGM gas temperature accompanying the end of helium reionization \\citep[see also][]{theuns2002c,fg2008a}.  However, the reality of this feature remains unclear.  Although a similar dip was found at modest significance by \\citet{dallaglio2008} in a sample of 40 high-resolution spectra, neither \\citet{mcdonald2005} nor \\citet{dallaglio2009} detected such a feature using samples of moderate-resolution SDSS spectra larger than that of  \\citet{bernardi2003}.  Indeed, \\citet{fg2008a} emphasized that the significance of their feature may have been exaggerated by the assumption of an underlying power law evolution in \\taueff.  The association of such a sharp feature with helium reionization has also been challenged on theoretical grounds.  \\citet{bolton2009b} argued that even if the IGM is photo-heated during the course of a very rapid ($\\Delta z \\simeq 0.2$) helium reionization, the expected signature would be a sudden decrease in \\taueff\\ followed by a gradual recovery driven by adiabatic cooling in the voids.  Moreover, recent temperature measurements indicate that heating due to helium reionization was an extended process with $\\Delta z \\gtrsim 1$ \\citep{becker2011a}, and so it is unclear why such a sharp feature in the opacity evolution should occur.  In this paper we present a new and highly precise measurement of the mean transmitted \\lya\\ flux over $2 < z < 5$ using a sample of 6065 moderate-resolution quasar spectra drawn from the seventh data release of the Sloan Digital Sky Survey (SDSS DR7).  The main innovation of this work is the introduction of a new method  that uses composite spectra to measure the overall shape the mean flux evolution to high accuracy without fitting continua in the \\lya\\ forest.  These results are then normalized to measurements made from high-resolution data at $z \\sim 2$, where continuum errors are minimal.  In common with all strategies that attempt to estimate quasar continua over the \\lya\\ forest based on redder, unabsorbed regions of the spectra (e.g., PCA analysis), it is essentially impossible to be completely certain that redshift-dependent variations in quasar spectral energy distributions (SEDs) are not affecting the derived $F(z)$ measurements.  We present tests, however, which demonstrate that such systematics are unlikely to be affecting our results at a level greater than the error bars. Our approach builds on the successful use of composite quasar spectra to statistically constrain other properties of the IGM, for example the mean free path of ionizing photons \\citep{prochaska2009b}.  We also take advantage of the improved statistical accuracy afforded by DR7 to increase substantially the number of quasars analyzed compared to previous works \\citep[][although \\paris\\ also used DR7 data]{bernardi2003,mcdonald2005,dallaglio2009,paris2011}.  In Section~\\ref{sec:data_method} we describe our technique, which uses composite spectra to compute \\meanfluxz\\ relative to that at $z \\sim 2$.  Our results are presented in Section~\\ref{sec:results}.  In Section~\\ref{sec:comparison}, we compare our measurements to results from the literature.  Finally, we summarize in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  We have presented new measurements of the mean transmitted flux in the \\lya\\ forest over $2.15 \\le z \\le 4.85$ made using composite quasar spectra.  Our sample includes 6065 individual spectra from the SDSS DR7 quasar catalogue, which were combined into 26 composites with mean quasar redshifts $z_{\\rm comp} = 2.25$ to 5.08.  The fact that the quasar composites all have highly similar fluxes redward of \\lya\\ suggests that the differences in the observed flux in the forest are driven by the evolution in the mean intergalactic \\lya\\ opacity.  The ratio of $F(z)$ at two different redshifts can be measured by comparing the observed mean flux in two composites at the same rest wavelength.  We used these ratios to directly constrain $F(z)$ as a fraction of $F(z=2.15)$.  We also conducted tests that address our assumption that the composites have similar unabsorbed continua over the \\lya\\ forest, finding that the errors in $F(z)/F(z=2.15)$ arising from variations in the continua are likely to be within  our uncertainty estimates.  The overall normalization for $F(z)$ was determined using literature values made from high-resolution data at $z \\le 2.5$ \\citep{fg2008a}, where continuum uncertainties, although accounted for, are relatively minor.  The statistical accuracy offered by the large number of quasar spectra means that this normalization factor now dominates the uncertainty in $F(z)$ at $z \\le 3.75$.  Our main results give the mean transmitted flux associated with systems of column density $N_{\\rm H\\,I} < 10^{19}~{\\rm cm^{-2}}$, and do not include the contribution from metal lines.  We fit the mean flux using two functions.  The first is a discreet $F(z)$ in bins of $\\Delta z = 0.1$, while the second is a modified power law of the form $\\tau_{\\rm eff}(z) = \\tau_0 \\left[ (1+z)/(1+z_{0}) \\right]^{\\beta} + C$, where $\\tau_{\\rm eff}(z) = -\\ln{F(z)}$.  Although previous works had generally parametrized \\taueffz\\ as a pure power law ($C=0$), we found that this was insufficient to describe the observed flux ratios at the required level of precision while simultaneously producing the correct normalization.  We emphasize that the true functional form of \\taueffz\\ is unknown, and that our parametrization is appropriate only in that it provides a reasonable match to the discreet $F(z)$ values.  Caution should be used when extrapolating this (or any) fit to higher and lower redshifts.  We also stress that, due to the fact that our results are generated from the ratios of fluxes at different redshifts, the parameters for both the discreet $F(z)$ and modified power law \\taueffz\\ are correlated, and that the full correlation matrices should be considered when applying our results.  The mean transmitted flux is found to evolve smoothly with redshift, which is consistent with a gradual evolution of the \\hi\\ ionization rate and gas temperature over $2 < z < 5$.  We see no evidence of previously reported features at $z \\simeq 3.2$, which appear to have been spurious.  Our results are otherwise generally consistent with the most careful measurements made to date using high-resolution data \\citep{fg2008a}, but offer a substantially greater level of precision and extend out to higher redshifts.  Our measurements also offer an improvement over previous efforts using SDSS spectra, which appear generally to have been affected by systematic errors \\citep[although see][]{paris2011}.  This work demonstrates the significant advantages of using large data sets to measure the mean transmitted flux.  In addition to the high degree of statistical accuracy, the use of composites makes it possible to avoid fitting continua at high redshifts, which is the primary source of systematic uncertainty in other approaches.  These measurements should therefore be useful to a variety of studies that use the \\lya\\ forest."
    },
    "1208/1208.5787.txt": {
        "abstract": "{ %Each chapter should be preceded by an abstract (10--15 lines long) that summarizes the content. The abstract will appear \\textit{online} at \\url{www.SpringerLink.com} and be available with unrestricted access. This allows unregistered users to read the abstract as a teaser for the complete chapter. As a general rule the abstracts will not appear in the printed version of your book unless it is the style of your particular book or that of the series to which your book belongs. %Please use the 'starred' version of the new Springer \\texttt{abstract} command for typesetting the text of the online abstracts (cf. source file of this chapter template \\texttt{abstract}) and include them with the source files of your manuscript. Use the plain \\texttt{abstract} command if the abstract is also to appear in the printed version of the book. %}  \\abstract{I outline, from a theoretical and somewhat personal perspective, significant features of Pulsar Wind Nebulae as Cosmic Accelerators.  I discuss recent studies of Pulsar Wind Nebulae (PWNe). I pay special attention to the recently discovered gamma ray ``flares'' in the Crab Nebula's emission, focusing on the possibility, raised by the observations, that the accelerating electric field exceeds the magnetic field, suggesting that reconnection in the persistent current layer (a ``current sheet'') plays a significant role in the behavior of this well studied Pevatron.  I address the present status of the termination shock model for the particle accelerator that converts the wind flow energy to the observed non thermal particle spectra, concluding that it has a number of major difficulties related to the transverse magnetic geometry of the shock wave. I discuss recent work on the inferred pair outflow rates, which are in excess of those predicted by existing theories of pair creation, and use those results to point out that the consequent mass loading of the wind reduces the wind's bulk flow 4-velocity to the point that dissipation of the magnetic field in a pulsar's wind upstream of the termination shock is restored to life as a viable model for the solution of the ``$\\sigma$'' problem.  I discuss some suggestions that current starvation in the current flow supporting the structured (``striped'') upstream magnetic field, perhaps inducing a transition to superluminal wave propagation. I show that current starvation probably does not occur, because those currents are carried in the current sheet separating there stripes rather than in the stripes themselves\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnote[2]{Collaborators, none of whom should be held responsible for the content of this paper: D. Alsop, E. Amato, D. Backer, P. Chang, N. Bucciantini, B. Gaensler, Y. Gallant, V. Kaspi, A.B. Langdon, C. Max, E. Quataert, A. Spitkovsky, M. Tavani,  A. Timokhin}}. \\renewcommand{\\thefootnote}{\\arabic{footnote}} %\\vspace*{1cm} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.3899_arXiv.txt": {
        "abstract": "We analyze age and metallicity estimates for an unprecedented database of some 5.5 million stars distributed throughout the Large Magellanic Cloud (LMC) main body, obtained from CCD Washington $CT_1$ photometry, reported on in Piatti et al. 2012. We produce a comprehensive field star Age-Metallicity Relationship (AMR) from the earliest epoch until $\\sim$ 1 Gyr ago. This AMR reveals that the LMC has not evolved chemically  as either a closed-box or bursting system, exclusively, but as a combination of both scenarios that have varied in relative strength over the lifetime of the galaxy, although the bursting model falls closer to the data in general. Furthermore, while old and metal-poor field stars have been preferentially formed in the outer disk,  younger and more metal-rich stars have mostly been formed in the inner disk, confirming an outside-in formation. We provide  evidence for the formation of stars between 5 and 12 Gyr, during the cluster age gap, although chemical enrichment during this period was minimal. We find no significant metallicity gradient in the LMC. We also find that the range in the metallicity of an LMC field has varied during the lifetime of the LMC. In particular, we find only a small range     of the metal abundance in the outer disk fields, whereas an average range of $\\Delta$[Fe/H] = +0.3 $\\pm$ 0.1 dex appears  in the inner disk fields. Finally, the cluster and field AMRs show a satisfactory match only for the last 3 Gyr,  while for the oldest ages ($>$ 11 Gyr) the cluster AMR is a remarkable lower envelope to the field AMR. Such a difference may be due to the very rapid early chemical evolution and lack of observed field stars in this regime, whereas the globular clusters are easily studied. This large difference is not easy to explain as coming from stripped ancient Small Magellanic Cloud (SMC) clusters, although the field SMC AMR is on average $\\sim$ 0.4 dex more metal-poor at all ages than that of the LMC but otherwise very similar. ",
        "introduction": "The Large Magellanic Cloud (LMC) has long been recognized as a fundamental benchmark for a wide variety of astrophysical studies. As the closest bulge-less dwarf disk galaxy \\citep{b12}, it has turned out to be the ideal local analog for the detailed study of these most common and primeval galaxies. Ages and abundances of LMC field star populations are prime indicators of the galaxy's chemical evolution and star formation history (SFH). This becomes even more relevant since its formation and chemical evolution cannot be fully traced from its star cluster populations, due to the well-known extended age gap. The LMC age-metallicity relationship (AMR) has been the subject of a number of studies \\cite[among others]{oetal91,hetal99,cetal05,retal11}. Among them, two perhaps best summarize our current knowledge in this field. First, \\citet[hereafter CGAH; see also references therein]{cetal11} have examined the AMR for field star populations, based on Calcium triplet spectroscopy of individual red giants and BVRI photometry in ten 34$\\arcmin$$\\times$33$\\arcmin$ LMC fields. They found that: i) the AMRs for their fields are statistically indistinguishable; ii) the disk AMR is similar to that of the LMC star clusters and is well reproduced by closed-box models or models with a small degree of outflow; iii) the lack of clusters with ages between 3 and 10 Gyr is not observed in the field population; iv) the age of the youngest population observed in each field increases with galactocentric distance; v) the rapid chemical enrichment observed in the last few Gyrs is only observed in fields with R$<$7kpc; vi) the metallicity gradient observed in the outer disk can be explained by an increase in the age of the youngest stars and a concomitant decrease in their metallicity; and vii) they find much better evidence for an outside-in than inside-out formation scenario, in contradiction to generic $\\Lambda CDM$ models. Secondly,  \\citet[hereafter HZ09; see also references therein)]{hz09} presented the first-ever global, spatially-resolved reconstruction of the SFH, based on the application of their StarFISH analysis software to the multiband photometry of twenty million stars from the Magellanic Clouds Photometric Survey. They found that there existed a long relatively quiescent epoch (from $\\sim$ 12 to 5 Gyr ago) during which the star formation was suppressed throughout the LMC; the metallicity also remained stagnant during this period. They concluded that the field and cluster star formation modes have been tightly coupled throughout the LMC's history.  Although these studies represent the state-of-the-art of our knowledge of the  LMC AMR, they leave unanswered a number of outstanding questions: What caused the general lull in  SF between $\\sim$ 5 and 12 Gyr ago? Are the cluster and field AMRs really tightly coupled? Can the LMC AMR best be described by a closed-box, bursting or other chemical evolution model? What, if any, are the radial dependences? In addition, HZ09 did not go deep enough to derive the full SFH from information on the Main Sequence (MS). They reached a limiting magnitude between $V$ = 20 and  21 mag, depending on the local degree of crowding in the images, corresponding to stars younger than 3 Gyr old on the MS if the theoretical isochrones of \\citet{getal02} and a LMC distance modulus of 18.5 mag are used. Thus, the advantages of covering an enormous extension of the LMC is partially offset by the loss in depth of the limiting magnitude. On the other hand, the ten fields of CGAH cover a rather small fraction of the whole LMC. Therefore, it is desirable to obtain an overall deeper AMR for the LMC which also covers a larger area. Previous AMRs have been founded on theoretical isochrones, numerical models, or synthetic Color-Magnitude Diagrams (CMDs), so that an AMR built from actual measured ages and metallicities is very valuable. A comprehensive comparison between the field and cluster AMRs obtained using the same procedure is also lacking. All these aims demand the availability of a huge volume of high quality data as well as a powerful technique to provide both accurate ages and metallicities.  In this paper we address these issues for the first time. We make use of an unprecedented database of some 5.5 million stars measured with the Washington $CT_1$ photometric system, which are spread over a large part of the LMC main body. From this database, we produce the LMC field AMR from the birth of the galaxy until $\\sim$ 1 Gyr ago, using the $\\delta$$T_1$ index and the standard  giant branch isoabundance curves to estimate ages and metallicities, respectively, of the most representative field populations. These provide approximately independent measurements of these two quantities, minimizing the age-metallicity degeneracy problem. In addition, this is the first overall LMC field star AMR obtained from Washington data; thus complementing those derived from other data sets such as HZ09 or the AMR obtained from Washington data for LMC clusters \\citep{p11a}. Finally, we homogeneously compared the derived field star AMR to that for the LMC cluster population with ages and metallicities put on the same scales using these two Washington datasets. This kind of comparison has not been accomplished before. The paper is organized as follows: Section 2 briefly describes the data handling and analysis from which \\citet{pgm12} estimated  the field star ages and metallicities. Section 3 deals with the aforementioned issue of a comprehensive AMR of the LMC field star population. In Section 4 we discuss our results and compare them with previous studies, while Section 5 summarizes our major findings. ",
        "conclusions": "In Fig. 5, we have overplotted with solid lines different field star AMRs along with our presently derived composite AMR, namely: HZ09 (yellow),  \\citet{retal11} (black), \\citet[hereafter PT98]{pt98} (blue), and  \\citet{getal98} (red). The red line AMR is based on a closed-box model, while the blue curve is a bursting model. We also included with red and blue filled circles the AMRs derived by \\citet{cetal08} for the LMC bar and disk, respectively. At first glance, we find that the bursting SFH modeled by PT98 appears to be the one which best resembles the AMR derived by \\citet{cetal08}, instead of closed-box models as Carrera et al. suggested. However, such a resemblance is only apparent since PT98 constructed their model using nearly no star formation from $\\sim$ 12 up to 3 Gyr ago (see their Fig. 2). This clearly contradicts not only Carrera et al.'s result but also ours, which show that there were many stars formed in the LMC in that period (see Fig. 4). Indeed, we actually see no significant chemical evolution from about 12 - 6 Gyr, even though stars were formed. In turn, the closed-box models appear to be qualitatively closer to HZ09's reconstructed AMR.  Since HZ09's AMR is based on a relatively bright limiting magnitude database and CGAH's AMRs rely on ages and metallicities for stars distributed in ten fields (each only slightly smaller than ours), we believe that the present composite AMR has several important advantages over these previous ones, and possibly reconciles previous conclusions about the major enrichment processes that have dominated the chemical evolution of the LMC from its birth until $\\sim$ 1 Gyr ago. Note that a large number of fields distributed through the galaxy are analysed here and their representative oldest MS TOs are well measured in all fields. The composite AMR we derive results in a complex function having HZ09's AMR (or alternatively the closed-box model) and CGAH's AMRs (or alternatively the bursting model) as lower and approximately upper envelopes in metallicity, respectively, although the bursting model is a much better fit. Therefore, we find evidence that the LMC  has not chemically evolved as a closed-box or bursting system, exclusively, but as a combination of both scenarios that likely have varied in importance during the lifetime of the galaxy, but with the bursting model dominating. The closed-box model presumably reproduces the metallicity trend that the LMC would have had if bursting formation episodes had not taken place. However, since the LMC would appear to have experienced such an enhanced formation event(s), important chemical enrichment has occurred from non well-mixed gas spread through the LMC. CGAH  also found that the AMRs for their ten fields are statistically indistinguishable. We note, however, that six of their fields are aligned somewhat perpendicular to the LMC bar, reaching quite low density outer regions, and therefore, that their coverage represents a relatively small percentage of the whole field population. We show in Fig. 1 that, when more field stars distributed through the LMC are analyzed with age and metallicity uncertainties robustly considered, distinct individual AMRs do arise. Indeed, Figs. 2 and 3 illustrate how different AMRs are for inner and outer fields.  When inspecting in detail our composite field LMC AMR, the  relatively quiescent epoch ($t$ $\\sim$ 5 to 12 Gyr) claimed by HZ09 and also frequently considered as a feature engraved in the cluster formation processes, i.e. the cluster age-gap \\cite[among others]{getal97,petal02,betal04} is not observed. On the contrary, there exists a noticeable number of fields with representative ages spanning the age gap (from $\\sim$ 12 Gyr to 3 Gyr), which further strengthens the difference between cluster and field star formation during this epoch. Of course, we do not quantitatively compare the level of SF in different epochs, we simply measure the properties of the dominant population. However, during this extended period, although some star formation occurred, it was not accompanied by any significant chemical evolution until starting $\\sim$6 Gyr ago. Again curiously, there were several Gyr of star formation and chemical evolution before the cluster age gap ended. In addition, although the ages estimated by CGAH  of field stars spanning the cluster Age Gap could have uncertainties necessarily large for individual stars, and consequently their SFH would still indicate a relatively  quiescent epoch between 5 and 10 Gyr as HZ09 pointed out, we provide here evidence of the existence of stars formed between 5 and 12 Gyr  that  represent the most numerous populations in their respective regions. Note that our metallicities are generally about 0.1 - 0.2 dex lower than CGAH for younger ages but higher for the oldest stars, indicating a smaller total chemical enrichment over the lifetime of the galaxy compared to that found by CGAH. Our agreement with \\citet{retal11} is somewhat better. We also find that the amount of chemical evolution (as measured by the increase   in the metallicity) of the LMC fields has varied during the lifetime of the LMC. Particularly, we find only a small range     of the metal abundance within the considered uncertainties for the outer disk fields, whereas an average increase of $\\Delta$[Fe/H] = 0.3 $\\pm$ 0.1 dex appears  in the inner disk fields, and this increase occurred over a relatively shorter time period. Hence, a bursting star formation scenario turns out to be a plausible explanation if the enhanced star formation is accompanied by a vigorous nucleosynthesis process that takes place during the burst.   Finally, we present a homogeneous comparison between the composite field AMR with that for 81 LMC clusters with ages ($\\ga$ 1 Gyr) and metallicities derived on the same scales as here. We use the ages and metallicities compiled by \\citet{petal11a} for 45 clusters observed in the Washington system, to which we add 36 clusters with ages estimated by \\citet{p11c} from similar data. We estimate here their metallicities following the same procedure used for the studied fields (see Section 2). The resulting cluster AMR is depicted in Fig. 6 with dark-gray filled boxes superimposed onto the composite field LMC (open boxes with error bars). As can be seen, the cluster AMR satisfactorily matches the field AMR only for the last 3 Gyr, while it is a remarkable lower envelope of the field AMR  for older ages ($t$ $>$ 11 Gyr). The most likely explanation is a very rapid early chemical enrichment traced by the very visible globular clusters, but their coeval, low metallicity field counterparts are so rare that they are missed in our data. The origin of the 15 oldest LMC clusters still remains unexplained and constitutes one of the most intriguing enigmas in our understanding of the LMC formation and evolution. Different studies show that they have very similar properties to the globular clusters in the Milky Way \\citep[][among others]{betal96,mg04,vdbm04,metal09,metal10}, except for their orbits, which are within the LMC disk instead of in an isothermal halo \\citep{b07}. On the other hand, Fig. 4 show that there exist field star populations older than 10 Gyr and about as old as the old globulars. These results go along with the curious conundrum of the absence of clusters during the infamous Age Gap \\citep{betal04}. Since HZ09 found that there was a  relatively quiescent epoch in the field star formation from approximately 12 to 5 Gyr ago (similar to that observed for star clusters), they also concluded that field and cluster star formation modes are tightly coupled. Notice that the ages and metallicities used by HZ09 for the 85 clusters are not themselves on a homogeneous scale nor on the same field age/metallicity scales.  In order to look for clues for the very low metallicities of the oldest LMC clusters, we reconstructed the cluster and field AMRs of the SMC, also from Washington photometry obtained by us. As for the field AMR we used the ages and metallicities derived by \\citet[his Table 4)]{p12} and applied to them the same binning and error analyses as for the composite LMC field AMR (Fig. 4). Note that these ages and metallicities are all set on the present age/metallicity scales. We also compiled 59 SMC clusters ($t$ $\\ge$ 1 Gyr) from \\citet{petal11b}, and \\citet{p11a,p11b} with ages and metallicities tied to the same scales. Fig. 6 shows the resulting SMC AMRs depicted with open triangles for its field stars and with filled triangles for its star cluster population. As can be seen, cluster and field stars apparently share similar chemical enrichment histories in the SMC, although the population of old clusters drastically decreases beyond $\\sim$ 7 Gyr and there is only 1 older than 10 Gyr. \\citet{p11b} showed, based on the statistics of catalogued and studied clusters, that a total of only seven relatively old/old clusters remain to be studied, and an even smaller number is obtained if the cluster spatial distribution is considered. From this result, we conclude that the SMC cluster AMR is relatively well-known, particularly towards its older and more metal-poor end. Therefore, it does not seem easy to connect the origin of the oldest LMC cluster population to stripping events of ancient SMC star clusters. Moreover, the composite SMC field AMR is on average $\\sim$ 0.4 dex more metal-poor at all ages than that of the counterpart in the LMC, with little variation, indicating that the global chemical evolution in these two galaxies was quite similar in nature but with an offset to lower metallicity in the SMC. In particular, there was a very early and rapid period of enrichment, followed by a long quiescent epoch with some star formation in both Clouds but cluster formation only in the SMC and little to no metallicity increase and finally a recent period of substantial enrichment starting about 6Gyr ago. This is in very good agreement with the SMC AMR found by \\citet{petal10}. The relative deficiency in heavy elements of the SMC could explain the metallicity of a few old LMC clusters, if they were captured from the SMC \\citep{betal12}, but this is an unlikely argument to explain the majority of them. In fact, it is curious in this context that the the oldest SMC cluster is at the young and metal-rich extreme of the LMC globular cluster distributions."
    },
    "1208/1208.1314_arXiv.txt": {
        "abstract": "We study the hydrostatic equilibrium structure of compact stars in the Eddington-inspired Born-Infeld gravity recently proposed by Ba\\~{n}ados and Ferreira~[Phys. Rev. Lett. {\\bf 105}, 011101 (2010)]. We also develop a framework to study the radial perturbations and stability of compact stars in this theory. We find that the standard results of stellar stability still hold in this theory. The frequency square of the fundamental oscillation mode vanishes for the maximum-mass stellar configuration. The dependence of the oscillation mode frequencies on the coupling parameter $\\kappa$ of the theory is also investigated. We find that the fundamental mode is insensitive to the value of $\\kappa$, while higher-order modes depend more strongly on $\\kappa$. ",
        "introduction": "\\label{sec:Intro}   General relativity (GR) has been the most successful and popular theory of gravity in the past century. From its classic predictions on the perihelion advance of Mercury and deflection of light, to the later predictions such as the orbital decay of the Hulse-Taylor binary pulsar due to gravitational-wave damping, GR has passed the experimental observations in these weak-field situations with flying colors. Testing the predictions of GR in strong-field situations, such as the final stage of binary black hole coalescence, will come soon from the detection of gravitational waves (see~Ref. \\cite{Will2006} for a review on experimental tests of GR).  Despite the great success of GR, the idea that GR may not be the correct theory to describe the Universe on cosmological scales is also gaining attention recently. This is due to the fact that, if GR is correct, we must then require the Universe to be dominated by some unknown component, called dark energy, in order to explain the accelerating expansion of the Universe. In the past decade, various alternative theories of gravity which deviate from GR on cosmological scales have been proposed in order to explain cosmological observations (see Ref. \\cite{Clifton2012} for a recent review).   On the other hand, it is also well known that GR is not complete because of its prediction of spacetime singularities in the big bang and those inside black holes. It is generally believed that a consistent theory of quantum gravity is needed to resolve this issue. Recently, a new Eddington-inspired Born-Infeld (EiBI) theory of gravity was proposed by Ba\\~{n}ados and Ferreira \\cite{Ferreira10}. The appealing properties of this theory are that it is equivalent to GR in vacuum and can avoid the big bang singularity \\cite{Ferreira10}. The theory differs from GR only in the presence of matter, and in particular the deviation becomes significant at high densities. It is thus reasonable to expect that compact stars may be the best astrophysical laboratories to test EiBI gravity. In fact, Pani {\\it et al.}~\\cite{Pani11,Pani12} have recently studied the structure of compact stars in EiBI gravity.  In this work, we report our investigation of compact stars in EiBI gravity. We first extend the work of Pani {\\it et al.}~\\cite{Pani11,Pani12} by studying the static structure of compact stars in this theory using a large set of realistic equations of state (EOS). Furthermore, we develop a framework to study the radial perturbations and stability of these stars in this theory for the first time.   The plan of the paper is as follows: In Sec. ~\\ref{sec:EiBI} we briefly summarize EiBI gravity. In Sec.~\\ref{sec:Static} we derive the equations for constructing static equilibrium stars in EiBI gravity. In Sec.~\\ref{sec:perturbation} we present the linearized equations for radial oscillations of compact stars. Section ~\\ref{sec:eos} presents the technique we use to employ realistic EOS models in this study. In Sec.~\\ref{sec:results} we present our numerical results. Finally, our conclusions are summarized in Sec.~\\ref{sec:conclusion}. We use units where $G=c=1$ unless otherwise noted. ",
        "conclusions": "\\label{sec:conclusion}   In this paper, we have studied the equilibrium structure of compact stars in EiBI gravity using a set of eight realistic EOS models. Our formulation of the structure equations is different from the recent works of Pani {\\it et al.}~\\cite{Pani11,Pani12} in such a way that the resulting differential equations resemble more closely the corresponding equations in GR. We solved the structure equations numerically and found that the maximum mass of a neutron star can be larger than that in GR when the parameter $\\kappa$ in EiBI gravity is positive. This implies that softer EOS models, which are ruled out in GR by the recent discovery of a neutron star with mass nearly $2 M_\\odot$, would be revived in EiBI gravity~\\cite{Pani12}.  We have also developed a theory of radial perturbation and studied the stability of compact stars in EiBI gravity by calculating the oscillation mode frequency square $\\omega^2$ for the first time. In contrast to the situation in GR, since the oscillation equation in EiBI gravity involves the second derivative $dP^2/ d\\epsilon^2$, we found that using standard techniques such as numerical interpolation or piecewise polytropic representation \\cite{Read09} to handle realistic EOS tables in the calculation would not produce reliable numerical results. This is because the data points in standard EOS tables are usually not dense enough to allow an accurate calculation of $dP^2/ d\\epsilon^2$. We thus followed Ref. \\cite{Haensel04} and used an 18-parameter analytic representation to model the EOS in our calculation. It should, however, be emphasized that the analytic fitting is not fundamentally essential to our stability analysis. It is employed in this work in order to improve numerical accuracy only.   We find that the standard results of stellar stability still hold in EiBI gravity. For a sequence of stars modeled by the same EOS, we found that the fundamental mode frequency square passes through zero at the central density corresponding to the maximum-mass configuration. Similar to the analysis of compact stars in GR, this point marks the onset of instability. Stellar models with central densities less than the critical point are stable because $\\omega^2 > 0 $. Furthermore, we also found that the criterion $dM/d\\epsilon_c > 0 $ does not guarantee stellar stability.   We have also studied the effects of the parameter $\\kappa$ on the oscillation modes. For a fixed stellar mass, we found that the fundamental mode frequency is insensitive to the value of $\\kappa$. On the contrary, the frequencies of higher-order modes depend strongly on $\\kappa$. In particular, a positive (negative) value of $\\kappa$ would decrease (increase) the frequencies of higher-order modes. Our results thus suggest that the detection of higher-order radial oscillation modes might provide useful constraints on the value of $\\kappa$ if they could be excited to large amplitudes. Of course, in reality oscillation modes could in general be excited to large amplitudes only in some catastrophic situations such as core collapse supernovae. However, those catastrophic events are in general highly nonspherical, and hence nonradial oscillation modes could become more relevant. Studying nonradial oscillations of compact stars and the corresponding gravitational wave signals would be a natural extension of our present work. We hope to return to this issue in the future and study whether the gravitational wave signals emitted by compact stars could be used to constrain EiBI gravity.     \\begin{appendix}"
    },
    "1208/1208.3311_arXiv.txt": {
        "abstract": "This paper builds on a calibration of the SNIa absolute distance scale begun with a core of distances based on the correlation between galaxy rotation rates and optical $I_C$ band photometry.  This new work extends the calibration through the use of mid-infrared photometry acquired at $3.6 \\mu$m with {\\it Spitzer Space Telescope}.  The great virtue of the satellite observations is constancy of the photometry at a level better than 1\\% across the sky.  The new calibration is based on 39 individual galaxies and 8 clusters that have been the sites of well observed SNIa.  The new $3.6 \\mu$m calibration is not yet as extensively based as the $I_C$ band calibration but is already sufficient to justify a preliminary report.  Distances based on the mid-infrared photometry are $2\\%$ greater in the mean than reported at $I_C$ band.  This difference is only marginally significant.  The $I_C$ band result is confirmed with only a small adjustment.  Incorporating a 1\\% decrease in the LMC distance, the present study indicates H$_0 = 75.2 \\pm3.0$~\\kmsMpc. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.1408_arXiv.txt": {
        "abstract": "{The environment appears to have a strong influence on fundamental properties of galaxies, modifying their morphology  and their  star formation histories.  Similarly, galaxies play a role in determining the properties of the hot intergalactic medium in groups, heating  and enriching it through a variety of mechanisms. NGC 5238 and NGC 4756 are the brightest unperturbed elliptical galaxies in their respective loose groups, but analysis of their environment suggest that they might be at different evolutionary stages.} {In the present study we aim at characterizing the properties of  the hot gas in the halos of the brightest members and in the environment.  In NGC 4756  we are also interested in the properties of a substructure identified to the SW and the region connecting the two structures, to search for a physical connection between the two.  However, we have to take into account the fact that the group is projected against the bright, X-ray emitting cluster A1361, which heavily contaminates and confuses the emission from the foreground structure.} {We present XMM-{\\it Newton} observations of the groups and the careful analysis to separate different components. We examine the X-ray morphology, hot gas distribution and  spectral characteristics of NGC 4756 and NGC 5328 and their companion galaxies. To better characterize the ambient, we also present a re-evaluation of the dynamical properties of the systems. SPH simulations are used to interpret the results.} {We find that the X-ray source associated with NGC 4756 indeed sits on top of extended emission from the background cluster A1361, but can be relatively well distinguished from it as a significant excess over it  out to $r \\rm \\sim150''\\ (\\sim 40\\ kpc)$.    NGC~4756 has an X-ray luminosity of  $\\rm  L_x\\sim 10^{41}$ erg s$^{-1}$  due to hot gas, with an average temperature of kT$\\sim 0.7 $ keV. We measure a faint diffuse emission also in the region of the subclump to the SW, but more interestingly, we detect gas between the two structures, indicating a possible physical connection. The X-ray emission from NGC 5328 is clearly peaked on the galaxy, also at $\\rm  L_x\\sim 10^{41}$ erg s$^{-1}$, and extends to r $\\sim 110$ kpc. Simulations provide an excellent reproduction of the SED and the global properties of both galaxies, which are caught at two different epochs of the same evolutionary process, with NGC 5328 $\\sim 2.5$ Gyr younger than NGC 4756. } {} ",
        "introduction": "Observationally, we see that a large fraction ($\\sim$50-60 \\%) of the galaxies in the Local Universe is in groups.  These could be isolated systems, or could be part of  filaments or chains, and  near/falling into clusters \\citep[see e.g.][]{Eke04, Tago08}. Groups, for which the galaxy velocity dispersion is comparable to the internal  velocities in galaxies,  provide a  controlled environment in which interactions (mergers)  act to modify galaxy properties, for example by removing the gas that fuels star formation.  Recent surveys have indeed verified that the transformation from star-forming systems into  ``passive'' ones  \\citep[][and references therein]{bai} is a result of the environmental action over the past 8 billion years, since redshift $z\\sim1$ \\citep{Kovac10,Peng10}. Cosmological hydrodynamical simulations by e.g. \\cite{Kobayashi05} and \\cite{Feldmann10, Feldmann11} to study the evolution of groups suggest that the  star-forming massive galaxies at the  center of group-type potentials at z$>>1 $ can become the massive, gas-poor early-type systems observed at the center of groups today.   \\begin{figure*} \\includegraphics[width=\\hsize,clip=true]{fig1} \\caption{ Top Left: Heliocentric radial velocities in a 1.5 Mpc box centered on NGC 4756. The mean heliocentric radial velocity and the 3$\\sigma$ velocity dispersion of the group are identified (dashed lines). The second  peak is due to A1631. Top right: B magnitude distribution of the group members. Bottom panels:  spatial distribution of galaxies centered on NGC 4756 within 1.5 Mpc, and a zoom corresponding to the  XMM-{\\it Newton} field of view (right), superposed on the 2D binned kernel-smoothed number density contours. Filled red squares are group members listed in \\cite{GrutzbaN4756}. The dashed circle is centered on the center of mass of the group and identifies the virial radius. } \\label{1pdf} \\end{figure*}    We therefore can expect that galaxies that we observe in clusters today are likely to have experienced  pre-processing in groups at some point in their history, before the group  itself fell into the cluster formed within the same infalling halo. In their study of close pairs in the Sloan Digital Sky Survey (SDSS), \\cite{Perez09} found  that galaxies are efficiently pre-processed by close encounters and mergers while in intermediate-density environments. With the same mechanism of interaction between pairs, mergers have probably depleted the reservoir of galaxies in the halo while building the dominant elliptical of the group. In addition to mergers, simulations suggest that ram-pressure stripping, once thought to operate only in rich environments, is active in the form of  ``strangulation'' in groups as well. This depletes gas rich galaxies from their hot ISM content, which is their largest gas reservoir, leaving  their molecular gas content nearly intact. In a time scale of about 1 Gyr, this process leads to quenching of star formation in gas rich disk galaxies, transforming them into S0s \\citep{Kawata08}. In fact  \\citet{Jeltema08} interpret the evidence that the $L_K - L_X$ relation for early-type galaxies in groups is systematically below that of field objects as an indirect evidence of hot gas stripped by viscous- or ram-pressure. This is supported by  a more direct evidence of the action of a stripping event seen in the X-ray image of the S0 galaxy NGC 6265, located $\\sim 250$ kpc west of the NGC6269 group (see also \\citealt{baldi}, or \\citealt{kim} for NGC 7619).  However, a different scenario is proposed by \\cite{Berrier}.  Their  simulations actually suggest that only a small fraction (30\\%) of galaxies today in cluster may have been accreted from groups, while the majority are accreted directly from the field and then reprocessed in the cluster environment.  To better understand processes related to galaxy evolution in different environments, a careful analysis of the dynamics and the characteristics of galaxies in groups in the local Universe provides the observational data necessary to interpret the different scenarios proposed.   \\cite{Tully10} shows that evolved groups, which are characterized by a high fraction of elliptical galaxies, have at least one  dominant E galaxy and relatively fewer intermediate luminosity galaxies  than pristine spiral rich groups \\citep[see e.g.][]{Grutzbauch09}. The evolution of massive Ellitpicals at the center of groups, expected to be the results of long-term evolution starting from very large scale, and the evolution of their hot gas content are beginning to be clarified through  self-consistent models \\citep{Bettoni12}.  In this context, we are using a multiwavelength approach to study groups (and their members) characterized by different observational properties (e.g., galaxy density and composition, dynamical properties) to understand whether these are representative of different evolutionary stages or different evolutionary paths.  We have obtained data at several wavelengths through IR, optical and UV observations with e.g., GALEX  and Spitzer \\citep{Marino10, Bettoni, annibali11, Marino12, panuzzo}. Here we present the results of XMM-Newton observations of two bright early-type galaxies in poor groups, NGC 4756 and NGC 5328, which we requested in order to study the properties of the hot gas in the galaxies and in the groups. Both groups should represent relatively late evolved stages, in which an  elliptical galaxy is the dominant, central member.   The hot gas in the central Elliptical should result both from stellar evolution (e.g. mass loss from evolved stars) and from the  environment \\citep{Mathews90,Mathews03}, and therefore is expected to be influenced by both the internal and the environmental  properties of the galaxy. Evolved groups with a dominant  early-type galaxy at their core typically show hot ($\\sim10^7$ K) halos detected in  X rays  \\citep[see e.g.][]{mul00,Mulchaey03}, but not all the details of the hot gas properties are clear at the present time.  We complement the X-ray data with a revision of the dynamical properties of the groups, and  with a set of smooth particle hydrodynamics (SPH) simulations which include chemo-photometric models.   These simulations  provide us with predictions for the properties of the system, in particular those of the  stellar population at different stages of evolution, to be compared with the observed  spectral energy distribution (SED), and of the hot gas (i.e. luminosity, mass..).  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{ \\includegraphics[clip=true]{fig2}} \\caption{Positions of the galaxies with measured velocities in the  NGC 4756 field, on the un-dithered V-band image of Abell 1631 obtained using WFI@ESO-MPI2.2m. Identified galaxies are at the velocity of NGC~4756 (red $\\sim 4000$ km s$^{-1}$) or at the velocity of A1631 (white, at $\\sim 13000$ km s$^{-1}$). Background galaxies are not labeled. } \\label{optical} \\label{vels} \\end{figure}    \\begin{figure*}  \\includegraphics[width=\\hsize]{fig3} \\caption{Same as  Fig.~\\ref{1pdf} for  NGC ~5328. Filled red squares are group members listed in \\citet{GrutzbaN5328} .} \\label{2pdf} \\end{figure*}    \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{fig4}} \\caption{XMM-Newton images of NGC 4756, separately for the three instrument. The positions of the detected sources are shown on the EPIC-M1 image. } \\label{n4756-fig1} \\end{figure*} ",
        "conclusions": "We have detected both galaxies with XMM-Newton data, and measured a sizable contribution from hot gas in both systems. The extent of the emission is larger than the optical size of the galaxies. We measure a total size of $\\sim$110 kpc for NGC 5328. NGC 4756 is observed against an X-ray emitting background cluster which prevents us from properly measuring the total extent of the source.  However we detect emission out to a radius of $\\sim$40 kpc around NGC 4756, and a connection to the SW compact-like group at $\\sim $130 kpc that is presumably interacting with NGC 4756 itself. The total luminosities measured are L$_x \\sim 10^{41}$ \\ergs, relatively low for groups, but this should be considered a lower limit for NGC 4756.  The results of SPH simulations, which describe the dynamical, morphological and chemo-photometric evolution in a self-consistent way,  suggests that NGC 4756 and NGC 5328 can be described by the same evolutionary path, stopped at two different epochs, with the last episode of star formations at 7 and 4.5 Gyr ago, respectively. The data and simulations sketch old and digested merging events at their origin, and they can now both be considered mature systems.  The predictions are in good agreement with observed quantities such as SED, $\\sigma$ for stars, and X-ray luminosity.  In spite of the apparent similarity of these two galaxies, in terms of luminosity, internal rotation, stellar content and evolutionary path, the evidence suggests that their respective groups are at different evolutionary phases.  A new analysis of the dynamics of both systems confirms that NGC 5328 sits at the center of a small group, which appears to be azimuthally symmetric and  relatively relaxed. NGC 5328 is the dominant galaxy in the system, about 2 mag. brighter than the next bright object.  The group has a relatively low virial mass, of a few $\\sim 10^{12}$ M$_\\odot$,  comparable to what is expected from a bright early-type galaxy.  The presence of  Abell 3574  at the same redshift suggests that this could be a satellite galaxy at the very periphery of the much larger potential, that has gone through a major preprocessing before being [eventually] acquired by the cluster.  NGC 4756 is composed of two main sub-condensations, NGC 4756 itself and the compact-like small group. The two are at the same average redshift, but might not have reached a dynamical equilibrium yet. The evidence of a connection in the X-ray emission between the two condensations suggests an interaction between the two structures which might eventually merge at a later stage of the evolution of the group."
    },
    "1208/1208.0829_arXiv.txt": {
        "abstract": "We present measurements of how multimode fiber focal-ratio degradation (FRD) and throughput vary with levels of fiber surface polish from 60 to 0.5 micron grit. Measurements used full-beam and laser injection methods at wavelengths between 0.4 and 0.8 microns on 17 meter lengths of Polymicro FBP 300 and 400\\mum\\ core fiber. Full-beam injection probed input focal-ratios between $f$/3 and $f$/13.5, while laser injection allowed us to isolate FRD at discrete injection angles up to 17 degrees ($f$/1.6 marginal ray). We find (1) FRD effects decrease as grit size decreases, with the largest gains in beam quality occurring at grit sizes above 5\\mum; (2) total throughput increases as grit size decreases, reaching 90\\% at \\filtI with the finest polishing levels; (3) total throughput is higher at redder wavelengths for coarser polishing grit, indicating surface-scattering as the primary source of loss. We also quantify the angular dependence of FRD as a function of polishing level. Our results indicate that a commonly adopted micro-bending model for FRD is a poor descriptor of the observed phenomenon. ",
        "introduction": "Multimode optical fibers provide the most cost-effective coupling between telescopes and spectrographs that allow spectrographs to be placed in stable environments. However, these fiber optics contribute to light loss from attenuation within the fiber material and surface-scattering of their ends, and increase entropy in the optical beam. The latter is referred to as focal ratio degradation (FRD), whereby light injected into a fiber at a particular \\fratio emerges at a faster (smaller) \\fratio. Ever since the first efforts to characterize FRD in astronomical applications\\cite{Angel77} astronomers have attempted to understand its cause(s) in the hope to lessen its effects\\cite{Carrasco,Oliveira}. Microbends have historically been a favored culprit\\cite{Carrasco,Gloge72}, but recently it has been suggested\\cite{Haynes11, Avila98} that scattering caused by surface-roughness on the fiber face contributes significantly to FRD.  We discuss the results of two experiments designed to measure how the amount of FRD depends on surface roughness, wavelength, and input angle. The experiments described here use both full-beam and laser injection methods\\cite{Carrasco} standard for FRD tests in astronomical applications. The full-beam method is useful for characterizing how a fiber would perform when fed by a telescope, and provides a straightforward way to compute practical metrics useful for designing spectroscopic instruments. The laser injection method allows light to be injected into the fiber at discrete input angles (compared to a filled ray-bundle cone). This angular dependence of scattering is a particularly sensitive diagnostic of the physical mechanisms responsible for FRD.  The method and results of our experiment of FRD dependence on surface roughness are presented in \\S \\ref{sec:polish}. Results for the wavelength dependence of FRD are reported in \\S \\ref{sec:wavelength}. The dependence of scattering on input angle is reported in \\S \\ref{sec:angle}, and the implications of our work are discussed in \\S \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} Two experiments to measure the amplitude of FRD as a function of wavelength, surface polish, and light-injection angle have been carried out and described. The primary results are: \\begin{enumerate}  \\item A component of FRD is attributable to the end-polish on fiber surfaces, however this appears to be a second-order effect relative to the impact of light-injection angle (beam speed) slower than $f$/3. FRD decreases with polishing down to finer grit sizes, but not significantly below grit-sizes of 5\\mum.  \\item Total throughput (light emerging at all angles) also depends on end-polish, with a wavelength dependence that indicates the increase in throughput is simply a reduction in surface-scattering.  The most significant gains occur for polishing that proceeds down to 5 $\\mu$m grit, although for most astronomical applications at low light-levels polishing down to the finest grit is measurably advantageous.  \\item The amount of FRD does \\textbf{not} depend on wavelength, as found now in several experiments. This is in contrast to some results and the predictions of micro-bend theory for FRD's origin\\cite{Carrasco}.  \\item FRD is dominated by light entering the fiber at smaller angles (in our case $<10^{\\circ}$), as is well known. Measurements here allow us to quantify this statement in detail. The amplitude and angular dependence of FRD also do not agree with predictions from micro-bend theory.  \\end{enumerate}  This work was supported by NSF grants ATI-0804576 and AST-1009471."
    },
    "1208/1208.5910_arXiv.txt": {
        "abstract": "We present spectroscopic analysis of the broad absorption line outflow in quasar SDSS J1512+1119. In particular, we  focus our attention on a kinematic component in which we identify \\pv\\ and \\siv/\\siv* absorption troughs. The shape of the unblended phosphorus doublet troughs and the three \\siv/\\siv* troughs allow us to obtain reliable column density measurements for these two ions. Photoionization modelling using these column densities and those of \\hei* constrain the abundance of phosphorus to the range of 0.5--4 times the solar value. The total column density, ionization parameter and metalicity inferred from the \\pv\\ and \\siv\\ column densities leads to large optical depth values for the common transition observed in BAL outflows. We show that the true \\civ\\ optical depth, is  $\\sim$~1000 times greater in the core of the absorption profile than the value deduced from it's apparent optical depth. ",
        "introduction": "\\label{intro}  Active galactic Nuclei (AGN) outflows have been detected as blueshifted broad absorption lines (BALs) in the UV spectra of $\\sim$ 20\\% of quasars \\citep{Hewett03,Dai08,Knigge08} and as narrow absorption lines (NALs) in $\\sim$ 50 \\% of Seyfert galaxies \\citep{Crenshaw03,Dunn08}. There is growing evidence that these ubiquitous sub-relativistic ionized outflows play an important role on sub-parsec as well as kilo-parsec scales in controling the growth of the central black hole, the evolution of the host galaxy and the chemical enrichment of the Intergatlactic medium (IGM) \\citep[e.g.][]{Elvis06,Moe09,Ostriker10}. Moreover, the study of the abundances of metals in these outflows observed over a range of redshifts (up to $z \\sim 5$) provide us with a unique probe to investigate the history and evolution of the chemical enrichment over cosmological scales, which constrains star formation scenarios and evolution of the host galaxy \\citep{Hamann97a,Hamann98,Hamann99,Hamann03,dimatteo04,Hamann07,Germain09,Barai11}.  Studying absorption lines from AGN outflows is the most direct way to determine chemical abundances in the AGN environment. This is done by comparing the column densities associated with ionized species of the different elements observed across the spectrum, combined with photoionization analysis. The major advantage of using absorption lines over emission lines in abundance studies resides in the fact that they provide diagnostics that largely do not depend on temperature and density \\citep{Hamann98}. Early abundance studies in BAL outflows implied particularly high abundances of heavy elements relative to hydrogen. In several cases, enhancement of carbon, nitrogen, oxygen and silicon by factors of tens to hundreds of times the solar values were reported in several objects \\citep[e.g.][]{Turnshek86,Turnshek96,Hamann98}, in contrast to the order of magnitude or less, generally derived from the analysis of the quasar emission lines \\citep[e.g.][]{Hamann93,Hamann02,Dietrich03,Juarez09}.  Perhaps the most puzzling observation was the detection of BALs associated with \\pv\\ \\citep{Junkkarinen95,Arav01,Hamann98,Hamann03}. Phosphorus is $\\sim$ 900 times less abundant than carbon in the solar photosphere \\citep{Lodders09}. Since \\pv\\ and \\civ\\ have similar ionization potentials % they are formed in similar environments. This suggests, based on direct comparison of the measured column densities, an overabundance of phosphorus over carbon of $\\gtsim$ 100 times the solar value \\citep[e.g.][]{Junkkarinen95,Junkkarinen97,Turnshek96,Hamann98}. \\citet{Shields96} suggested a scenario consistent with the reported phosphorus overabundances in which the enrichment of the BAL material is mainly controlled by a population of galactic novae. However, our group \\citep{Arav97,Arav99ak,Arav99,deKool01,Arav01,Arav01b,Arav02,Arav03,Scott04,Gabel05a} and others \\citep{Barlow97b,Hamann97b,Telfer98,Churchill99,Ganguly99} showed that column densities derived from the apparent optical depth analysis of BAL troughs are unreliable due to non-black saturation in the troughs. Therefore, \\citet{Hamann98,Hamann99,Hamann03}; and  \\citet{Leighly09} suggested that the extreme overabundance of phosphorous relative to carbon is an artifact of very high levels of saturation in the \\civ\\ troughs, compared to only mild (or non) saturation in the \\pv\\ troughs. Subsequent measurements of abundances in Seyfert and quasar outflows accounted for non-black saturation and yielded abundances of only a few times solar in the outflows \\citep{Gabel06,Arav07}.  The non-black saturation hypothesis was largely accepted by the community to explain the \\civ/\\pv\\ BAL observations. But this scenario implies that the actual optical depth in the \\civ\\ trough is roughly 1000 times larger than the apparent one, an assertion that was never verified empirically. In this paper, we study the UV outflow of SDSS J1512+1119, which exhibits deep absorption troughs from \\pv\\ as well as \\siv. In particular, we report the detection of the excited \\siv\\ $\\lambda$1073.51 line, a transition ten times weaker than the excited \\siv\\ $\\lambda$1072.96. Together with estimates of the number density provided by the analysis of absorption troughs from excited states of \\ciii\\ and \\feiii, we pinpoint the \\siv\\ column density. Photoionization modeling, using the derived column densities as input, shows that the phosphorus abundance is close to the solar values. This allows us to confirm that the true \\civ\\ optical depth is  $\\sim$~1000 times greater in the core of the absorption profile than the value deduced from apparent optical depth measurements.   The plan of the paper is as follows: In \\S~\\ref{dataredu} we present the VLT/X-Shooter observations of SDSS J1512+1119 along with the reduction of the data. In \\S~\\ref{anabs} we identify the spectral features and estimate the column density associated with each ionic species. We discuss the photoionization solution for the absorber and the implied phosphorus abundance in \\S~\\ref{discu}. We conclude the paper by summarizing the key points of the analysis in \\S~\\ref{conclu}. ",
        "conclusions": "\\label{conclu}  In this paper, we studied the UV outflow associated with the quasar SDSS J1512+1119 on the basis of new, medium resolution VLT/X-Shooter data. The extended wavelength coverage of the instrument allowed us to detect the outflow components in a multitude of ionic species. In particular we report the detection of deep \\pv\\ absorption troughs in kinematic component 2 as well as the detection of \\siv\\ and \\siv*. A detailed analysis of the \\siv* line profile allowed us to detect the weak $\\lambda 1073.51$ transition revealing a \\siv\\ column density larger than suggested by the apparent depth of the absorption troughs of the $\\lambda 1072.97$ transition.  Photoionization modeling of the absorber revealed that the absorber is thick, though the non-detection of significant \\feii\\ absorption troughs guarantees the absence of a significant amount of \\hi\\ bound-free opacity. Our accurate determination of the total \\pv\\, \\siv\\, \\hei* and lower ionization species column densities allowed us to characterize the physical state of the gas. We find that for the range of ionization parameters relevant for the present absorber, the phosphorus abundance relative to helium is consistent with the solar value. Using the parameter derived from the photoionization analysis, we show, as suggested in \\citet{Hamann98}, that a line such as the ubiquitous \\civ\\ is heavily saturated. The \\civ\\ column density derived from the apparent depth of the absorption line profile underestimates the column density by a factor of $\\sim$1000, providing a very poor estimate of its true column density.  The  phosphorus abundance we find is in disagreement with the extreme phosphorus abundances reported in the early literature \\citep[e.g.][]{Junkkarinen95,Junkkarinen97,Hamann98}. Other elemental abundances are found to be in agreement with the solar values. The fact that the abundances are similar to the solar values for an odd (\\pv) and even (\\siv) element points to enrichment by relatively ``normal'' galactic stellar populations \\citep[e.g.][]{Hamann97a} rather than the more exotic mechanism proposed by \\citet{Shields96} that would significantly enhance the P/S ratio \\citep{Hamann98}."
    },
    "1208/1208.2701_arXiv.txt": {
        "abstract": "Based on a suite of state-of-the-art high-resolution $N$-body simulations, we revisit the so-called halofit model (Smith et al. 2003) as an accurate fitting formula for the nonlinear matter power spectrum. While the halofit model has been frequently used as a standard cosmological tool to predict the nonlinear matter power spectrum in a universe dominated by cold dark matter, its precision has been limited by the low-resolution of $N$-body simulations used to determine the fitting parameters, suggesting the necessity of improved fitting formula at small scales for future cosmological studies. We run high-resolution $N$-body simulations for 16 cosmological models around the Wilkinson Microwave Anisotropy Probe (WMAP) best-fit cosmological parameters (1, 3, 5, and 7 year results), including dark energy models with a constant equation of state. The simulation results are used to re-calibrate the fitting parameters of the halofit model so as to reproduce small-scale power spectra of the $N$-body simulations, while keeping the precision at large scales. The revised fitting formula provides an accurate prediction of the nonlinear matter power spectrum in a wide range of wavenumber ($k \\leq 30h$\\,Mpc$^{-1}$) at redshifts $0 \\leq z \\leq 10$, with $5\\%$~precision for $k\\leq1\\,h$\\,Mpc$^{-1}$ at $0 \\leq z \\leq 10$ and $10\\%$ for $1 \\leq k\\leq 10\\, h$\\,Mpc$^{-1} $ at $0 \\leq z \\leq 3$. We discuss the impact of the improved halofit model on weak lensing power spectra and correlation functions, and show that the improved model better reproduces ray-tracing simulation results. ",
        "introduction": "The large-scale structure of the Universe has evolved under the influence of cosmic expansion and gravity, and its statistical nature contains valuable cosmological information. Among others, the power spectrum $P(k)$ is one of the most fundamental statistical quantities characterizing the large-scale structure. It has widely been used for cosmological studies, both in predicting various observable quantities and in extracting cosmological information from the observations~\\citep[e.g.,][]{pee93,dod03}. Given growing interests in high precision cosmological observations, of particular importance is an accurate theoretical template of the power spectrum, taking account of the nonlinear gravitational evolution.  Weak lensing induced by the large-scale structure between observed galaxies and the observer provides a unique opportunity to directly probe matter inhomogeneities in the Universe. This cosmic shear signal has been measured with a high signal-to-noise ratio by current large surveys including Canada-France-Hawaii Telescope Legacy Survey~\\citep[CFHTLS;][]{fu08}, Sloan Digital Sky Survey~\\citep[SDSS;][]{lin11,huff11}, and Cosmic Evolution Survey~\\citep[COSMOS;][]{massey07,sch10}. These surveys provided useful constraints on the cosmological parameters such as the matter density parameter $\\Omega_{\\rm m}$ and the amplitude of density fluctuation $\\sigma_8$. Future surveys such as Subaru Hyper Suprime-Cam~\\citep[HSC;][]{miyazaki06}, Dark Energy Survey~\\citep[DES;][]{des05}, and Large Synoptic Survey Telescope~\\citep[LSST;][]{lsst09} aim at measuring the cosmic shear signal with unprecedented precisions. While weak lensing probes matter fluctuations projected along the line-of-sight, one can extract the redshift evolution of the fluctuations, and hence accurate information on dark energy, using a technique called lensing tomography~\\citep[e.g.,][]{Hu99,TJ04} or a cross-correlation with intervening objects \\citep[e.g.,][]{oguri11}. However, accurate and unbiased cosmological constraints from these lensing measurements can be obtained only if we have appropriate likelihood function with given marginal distributions~\\citep{sit10,sit11} and an accurate model of the power spectrum $P(k)$. For instance, \\citet{ht05} argued that we typically need a few percent accuracy of $P(k)$ at the wavenumber $k<10h$\\,Mpc$^{-1}$ in order for the uncertainty of $P(k)$ not to degrade cosmological constraints in DES and LSST ~\\citep[see also][in which a similar conclusion is obtained]{eif11,hzm12}.  In the linear and quasi-linear regime of density fluctuations, the power spectrum can be computed for any given initial conditions and cosmological parameters using perturbation theory~\\citep[e.g.,][for a review]{ber02}. In the nonlinear regime, however, one has to resort to cosmological $N$-body simulations to study the nonlinear gravitational evolution. $N$-body simulation results are then used to develop phenomenological halo models or fitting formulae of nonlinear gravitational clustering. For instance, \\citet{pd96} provided a fitting formula of $P(k)$ based on a scaling ansatz presented in \\citet{hklm91}. \\citet[][hereafter S03]{sm03} proposed a new model of $P(k)$, the so-called halofit model, which is based on a halo model of structure formation \\citep[e.g.,][]{mf00,sel00,CS02}. In this halo model, all the matter content in the Universe is assumed to be bound in dark matter halos. Then the power spectrum is decomposed into two terms, the so-called one- and two-halo terms. The one-halo term describes matter correlations within the same dark matter halo, and is determined by the density profile of each halo. On the other hand, the two-halo term arises from the correlation between two distinct halos. The one-halo term  dominates at small scales, whereas the two-halo term dominate at large scales. The halofit model chose the functional form of $P(k)$ based on the halo model, but the model parameters were calibrated from $N$-body simulation results.  The halofit model by S03 is widely used to calculate the nonlinear matter power spectrum, yet it has been reported that the model fails to reproduce recent high-resolution $N$-body simulation results at small scales~\\citep[e.g.,][]{spr05,hil09,sato09,boy09,taka11,kie11,vn11a,vn11b,har12,cas12, it12}. For instance, \\citet{wv04} first pointed out that the halofit predicts a smaller power than their numerical results at small scales. \\citet{heit10} ran a suite of high-resolution simulations, called ``Coyote Universe'', and showed that $P(k)$ predicted by the halofit is $\\sim 5\\%$ smaller than their numerical results at $k \\sim 1h$\\,Mpc$^{-1}$. The one reason of the difference comes from the fact that the $N$-body simulations used in S03 have lower spatial resolution than latest ones. The another reason is that the halofit model in S03 is the fitting function for the Cold Dark Matter (CDM) model without baryons\\footnote{The presence of a significant fraction of baryon suppress the linear power spectrum at small scales. The fitting function in S03 is evaluated from the input linear power spectrum. Hence, the fitting function is slightly biased for the cosmological models with baryons.}. An outcome of the Coyote Universe simulations is a publicly available code ``cosmic emulator'' to calculate the nonlinear matter power spectrum by interpolating the simulations results for 38 different cosmological models \\citep{law10}. However, their emulator is restricted to a narrow range in $k<3h{\\rm Mpc}^{-1}$ and at low redshift $0 \\le z \\le 1$. Also, the Hubble parameter is automatically specified in the code using the cosmic microwave background (CMB) anisotropy constraint on the distance to the last scattering surface.  In this paper, we revisit the halofit model based on state-of-the-art high-resolution $N$-body simulations in 16 cosmological models around the Wilkinson Microwave Anisotropy Probe (WMAP) best-fit cosmological parameters. We allow the dark energy equation of state $w$ to deviate from $-1$, assuming that $w$ does not evolve with redshift. The halofit model has been tested for dark energy models ($w \\neq -1$) using N-body simulations ~\\citep[e.g.,][]{mtc06,ma07,cas09,fll09,ali10,cas11b}. While the original halofit model in S03 contains $30$ parameters, we increase the number of parameters to $35$ in order to achieve a better fit to the simulations. The new formula we present, which is summarized in Appendix, is widely applicable in the wavenumber range of $k<30h$\\,Mpc$^{-1}$ and the redshift range of $0 \\leq z \\leq 10$. Simply replacing the parameters in the original halofit model with new ones in the Appendix, an accuracy of fitting function is improved especially at small scales.  The present paper is organized as follows. In Section~\\ref{sec:nbody}, we begin by describing the $N$-body simulations and cosmological models used for the power spectrum analysis. Combining the $N$-body results with different box sizes, we discuss in detail the convergence of power spectrum measurement over the wide range of wave number. In Section~\\ref{sec:halo}, we re-calibrate the halofit model, and the revised version of the halofit model, whose explicit formula is given in Appendix, is compared with our $N$-body simulations. As an important implication of revised halofit model, in Section~\\ref{sec:wl}, we compute weak lensing power spectra, and compare them with direct ray-tracing simulation results, particularly focusing on the small-scale behavior. Finally, Section~\\ref{sec:conc} is devoted to conclusion and discussion. ",
        "conclusions": "\\label{sec:conc}  The halofit model presented in S03 has widely been used as a standard cosmological tool to predict the nonlinear matter power spectrum. However, it has been argued that the halofit model fails to reproduce recent high-resolution simulation results such that it underestimates the power spectrum by a few ten percent at small scales ($k \\gtrsim 0.1h$\\,Mpc$^{-1}$). The difference is crucial for analysis of upcoming weak lensing surveys such as Subaru HSC survey, DES, and LSST. In this paper, we have revisited the halofit model using the high resolution simulations for the $16$ cosmological models around the WMAP best-fit cosmological parameters, including the variation in the equation of state of dark energy.  The revised fitting formula can reproduce the simulation results very well in the range of $k<30h$\\,Mpc$^{-1}$ and $0 \\leq z \\leq 10$. Our new fitting formula is summarized in Appendix, which can easily be updated from the original halofit model by simply replacing parameters in original model with new values as well as adding a few terms. Our revised halofit is now implemented in current version of CAMB\\footnote{CAMB home page: http://camb.info/} (Oct. 2012), and hence one can easily calculate the nonlinear power spectrum $P(k)$ and the resulting weak lensing power spectra $C_\\ell$ and lensed CMB power spectrum $C_\\ell$ in our revised halofit model using CAMB.   We comment on effects of baryon cooling and massive neutrinos, both of which affect $P(k)$ at small scales. The baryon cooling would enhance the power at small scales by some ten percent at $k=10h$\\,Mpc$^{-1}$ and the enhancement becomes more significant for smaller scales~\\citep[e.g.,][]{jing06,rudd08,cas11,vd11,cas12}. However, the reliability of simulation results strongly relies on galaxy formation models they adopted. For example, \\cite{vd11} showed that the AGN feedback can decrease the power spectrum by $>10\\%$ at $k>1h$\\,Mpc$^{-1}$. The massive neutrinos also suppress the growth of density fluctuation below the so-called freestreaming scales~\\citep[e.g.,][]{bh09,bir12}. The power spectrum is suppressed by a few ten percent at small scales of $k \\simeq 0.1h$\\,Mpc$^{-1}$, depending on the total mass of neutrinos. Even though these effects can modify the small-scale nonlinear matter power spectrum, an accurate knowledge of the original (dark matter only) nonlinear power spectra as presented in this paper is still important as an ingredient for building models of more realistic nonlinear power spectra which take these effects into account.  Finally, while we have improved the fitting formula using the simulations, there are also several attempts to improve the halo model analytically. For instance, combining the perturbation theory at large scales with the halo model at small scales, \\cite{vn11a,vn11b} and \\cite{vsn12a,vsn12b} presented an improved halo model.  On smaller scales, \\cite{gio10} provided a prediction of the power spectrum using the halo model including the effect of substructure in the individual halo. These models also reproduce the simulation results well and are consistent with our fitting formula."
    },
    "1208/1208.2471_arXiv.txt": {
        "abstract": "Ultra High Cosmic Rays)  made by He-like lightest nuclei might  fit clustering along Cen A. Moreover He like UHECR nuclei explain Virgo  absence because the light nuclei fragility and opacity above a few Mpc. We foresaw (2009) that UHECR He from Cen-A AGN being fragile should partially fragment into secondaries  at  tens EeV  multiplet (D,$He^{3}$,p) as the recent twin multiplet discovered ones (AUGER-ICRC-2011), at $20$ EeV  along Cen A UHECR clustering.  We suggest that UHECR are also heavy radioactive galactic nuclei  as $Ni^{56}$,  $Ni^{57}$  and $Co^{57}$,$Co^{60}$ widely bent (tens degree up to $\\geq 100^{o}$)  by galactic fields.  UHECR radioactivity (in $\\beta$ and $\\gamma$ channels) and decay in flight at hundreds keV  is boosted (by huge Lorentz factor $\\Gamma_{Ni}\\simeq 10^{9}- 10^{8}$)  leading to PeVs electrons and consequent synchrotron TeVs gamma offering UHECR-TeV correlated sky anisotropy. Electron and tau  neutrinos secondaries at PeVs maybe the first signature of such expected radioactive secondary tail. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0368_arXiv.txt": {
        "abstract": "We report the discovery of a wide ($\\sim$1400~AU projected separation), common proper motion companion to the nearby M dwarf LHS~2803 (PSO~J207.0300$-$13.7422). This object was discovered during our census of the local T dwarf population using Pan-STARRS1 and 2MASS data. Using IRTF/SpeX near-infrared spectroscopy, we classify the secondary to be spectral type T5.5. University of Hawai`i 2.2m/SNIFS optical spectroscopy indicates that the primary  has a spectral type of M4.5, with approximately solar metallicity and no measurable H$\\alpha$ emission. We use this lack of activity to set a lower age limit for the system of 3.5 Gyr. Using a comparison with chance alignments of brown dwarfs and nearby stars, we conclude that the two objects are unlikely to be a chance association. The primary's photometric distance of 21 pc and its proper motion implies thin disk kinematics. Based on these kinematics and its metallicity, we set an upper age limit for the system of 10 Gyr. Evolutionary model calculations suggest that the secondary has a mass of 72$\\pm^{4}_{7}$ M$_{Jup}$, temperature of 1120$\\pm$80 K, and $\\log g$=5.4$\\pm$0.1 dex. Model atmosphere fitting to the near-IR spectrum gives similar physical parameters of 1100 K and $\\log g$=5.0. ",
        "introduction": "The past two decades have seen an explosion of interest in the study of ultracool dwarfs (spectral types of M6 or later), which can be stars or brown dwarfs. However, determining the physical properties of brown dwarfs, such as their mass, age and radius, is complicated due to their substellar nature. Unlike main-sequence stars, the lack of hydrogen burning means that brown dwarfs have no stable luminosity and hence there is a degeneracy between effective temperature, mass and age. This degeneracy can be broken for two classes of objects. The first type, \\citep[``mass benchmarks'';][]{Liu2008}, are binary brown dwarfs whose dynamical mass can be measured from their orbits. There are over a dozen such systems known \\citep{Dupuy2011}.  The second type, ``age benchmarks'', have their ages inferred from associated stars. This can be done by studying low-mass members of stellar clusters. However, the confirmed, young substellar populations of such associations are currently limited to mid-M to early-L dwarfs \\citep[e.g. ][]{Lodieu2012} and is hence not directly applicable to cooler field objects. Another source of benchmarks is substellar companions to higher mass stars. These objects have their ages determined from their stellar primaries and span a wide range in spectral types \\citep{Faherty2010}. Three close T dwarf companions have been discovered by high contrast imaging --- Gl~229B \\citep{Nakajima1995}, GJ~758B \\citep{Thalmann2010}, and SCR~1845$-$6357B \\citep{Biller2006} --- along with eleven wide companions to main sequence stars discovered using wide-field surveys or conventional imaging  --- GL~570D \\citep{Burgasser2000}, HD~3651B \\citep{Mugrauer2006,Luhman2007}, HN~Peg~B \\citep{Luhman2007}, Wolf~940B \\citep{Burningham2009}, Ross~458C \\citep{Goldman2010}, HIP~63510C \\citep{Scholz2010B}, HIP~73786B \\citep{Scholz2010B,Murray2011}, $\\epsilon$~Indi~Bab \\citep{Scholz2003, McCaughrean2004}, Gl~337CD \\citep{Wilson2001,Burgasser2005}, HIP~38939B \\citep{Deacon2012}, and BD~+01$^{\\circ}$~2920B \\citep{Pinfield2012}. Additionally \\cite{Albert2011} identify a potential companion to HD~15220 (CFBDS~J022644$-$062522) while \\cite{Dupuy2012} establish ULAS~1315$+$0826~\\citep{Pinfield2008} as a potential companion to TYC 884-383-1. Finally, two ultracool companions have been discovered to white dwarfs, LSPM 1459+0857B \\citep{Day-Jones2010} and WD~0806$-$661b \\citep{Luhman2011}, the latter of which has no spectroscopic data but is likely to be of the newly discovered Y spectral class~\\citep{Cushing2011}. These benchmarks can be used to test atmospheric and evolutionary models for mutual consistency by comparing their expected effective temperatures from their age and luminosity to the effective temperatures from atmospheric model fits to their spectra. The analysis of benchmark T dwarfs presented in \\cite{Deacon2012} indicates that the most recent atmospheric models \\citep{Allard2010,Saumon2008} match the effective temperatures inferred from evolutionary models to within $\\sim$100 K.  The vast majority of wide brown dwarf companions have been discovered by mining large astronomical surveys. One of the leading survey telescopes currently in operation is Pan-STARRS1 \\citep[PS1, ][]{Kaiser2002}. This 1.8m wide-field telescope sits atop Haleakal\\={a} on Maui in the Hawaiian Islands. The facility is operated by a consortium of US, German, UK and Taiwanese institutions and has been conducting full survey operations since May 2010. The telescope carries out a variety of astronomical surveys covering a wide range of science areas, from the solar system \\citep{Hsieh2012}, to the solar neighborhood \\citep{Deacon2011,Liu2011}, to high redshift supernovae \\citep[e.g. ][]{Chomiuk2012}. The most valuable of these surveys for ultracool dwarf science is the 3$\\pi$ Survey. This is covering the three-quarters of the sky ($\\delta > -30^{\\circ}$) visible from Maui in five filters in the Pan-STARRS1 system \\citep[$g_{P1}$, $r_{P1}$, $i_{P1}$, $z_{P1}$, $y_{P1}$, ][]{Tonry2012}. Two pairs of images will be taken per year in each passband for each area of sky during each of the three years of planned survey operation. Due to the fact that Pan-STARRS1 is only now finishing its second year of survey operation, our search for ultracool T dwarfs using PS1 has so far used data from the 2MASS survey \\citep{Skrutskie2006} as an additional epoch for proper motion determination and as a source of near-infrared photometry. This has resulted in the discovery of four relatively bright ($J \\sim 16.5$ mag) T dwarfs in PS1 commissioning data \\citep{Deacon2011} and, in combination with data from the Widefield Infrared Survey Explorer \\citep[WISE. ][]{Wright2010}, the identification of PSO~J043.5395+02.3995, a  nearby T8 brown dwarf with an extremely high proper motion (\\citealt{Liu2011} see also \\citealt{Scholz2011,Kirkpatrick2011}). Using the proper motion catalog produced for the PS1-2MASS T dwarf search, a dedicated search for cool companions to main sequence stars is also underway. This has already resulted in the discovery of HIP~38939B, a T4.5 companion to a mid-K star in the solar neighborhood \\citep{Deacon2012}.  Here we present the discovery of a mid-T-type companion\\footnote{This object was independently discovered by Muzic et al. AJ submitted\\nocite{Muzic2012}.} to the M dwarf LHS~2803 (PSO~J207.0300$-$13.7422). The primary is of a much lower mass than most other stars with a wide ($>$1000 AU) T dwarf companion. We use spectroscopic observations to characterize the both the primary and secondary stars. These data are then used to fit substellar evolutionary and atmospheric models to the secondary, allowing a comparison of the best fit models. ",
        "conclusions": "We have identified a common proper motion companion to the M dwarf LHS~2803. Spectroscopic observations classify the primary as M4.5$\\pm$0.5 and the secondary as T5.5$\\pm$0.5. Based on the literature photometric distance of 21$\\pm$3~pc, we calculate a projected separation of 1400$\\pm$200 AU. This system is among the widest known substellar companions to an M dwarf. We use the primary's lack of H$\\alpha$ emission to set a lower limit of 3.5 Gyr on the system age and its disk kinematics and approximately solar metallicity to set an upper age bound of 10 Gyr. Based on this age range, evolutionary model calculations indicate that the secondary has $T_\\mathrm{eff}$=1120~$\\pm$80~K and log~$g$=5.4~$\\pm$~0.1 with a mass of $72\\pm^{4}_{7}$~M$_{Jup}$, suggesting a relatively old, higher gravity object close to the maximum possible mass for a T dwarf. Comparing atmospheric models to our near-IR spectrum give a best fit of 1100~K and log~$g$=5.0, within one grid point of our evolutionary model calculations. The effective temperatures from the evolutionary and atmospheric models are in good agreement, in common with the similar benchmark mid-T dwarfs HN~Peg~B and HIP~38939~B \\citep{Deacon2012}. Pan-STARRS1 is ideally suited for identifying substellar companions to nearby stars, and a dedicated effort is underway to identify such objects."
    },
    "1208/1208.5251_arXiv.txt": {
        "abstract": "We prove that merging supermassive black holes (SMBHs) typically have neither equal masses, nor is their mass ratio too extreme. The majority of such mergers fall into the mass ratio range of $1:30$ to $1:3$, implying a spin flip during the inspiral. We also present a simple expression for the final spin $\\chi _{f}$ of the emerging SMBH, as function of the mass ratio, initial spin magnitudes, and orientation of the spins with respect to the orbital plane and each other. This formula approximates well more cumbersome expressions obtained from the fit with numerical simulations. By integrating over all equally likely orientations for precessing mergers we determine a lower approximant to the final spin distribution as function of the mass ratio alone. By folding this with the derived mass ratio dependent merger rate we derive a \\textit{lower} bound to the typical final spin value after mergers. We repeat the procedure deriving an \\textit{upper} bound for the typical spin in the case when the spins are aligned to the orbital angular momentum, such that there is no precession in the system. Both slopes of $% \\chi _{f}$ as function of the initial spins being smaller than one lead to two attractors at $\\chi _{f}^{prec}=0.2$ and $\\chi _{f}^{align}=0.45$, respectively. Real mergers, biased toward partial alignment by interactions with the environment (accretion, host galaxy, etc.) would generate a typical final spin lying between these two limiting values. These are the typical values of the spin after the merger, starting from which the spin can built up by further gaseous accretion. ",
        "introduction": "Einstein's general theory of relativity predicts that the coalescence of two compact objects (neutron stars or black holes) is accompanied (and driven by) intense gravitational radiation. Stellar mass (a few to a few ten solar masses, M$_{\\odot }$) black hole binaries emit gravitational waves with frequency falling into the best sensitivity range of LIGO \\cite{LIGO}, Virgo \\cite{Virgo} and GEO600 \\cite{GEO} Earth-based interferometric gravitational wave detectors. Up to date there are very few observations \\cite{IMBH} indicating the existence of intermediate mass black holes. Binaries formed by such black holes would emit gravitational waves falling into the frequency range of third generation gravitational wave detectors, like the Einstein Telescope \\cite{ET}.  Supermassive black holes (SMBHs) with masses of $10^{6}\\div 3\\times 10^{9}$ M% $_{\\odot }$ (or perhaps even higher) on the other hand are quite frequent, residing in the centre of each sufficiently massive galaxy. Their growth occurs by accretion phases and by mergers, the estimated contribution of each of these processes to the growth of mass being model-dependent. Recent observations \\cite{Kormendy} show that some galaxies never merge, and yet may have central back holes; representing either the pure SMBH birth population, or the birth population with some gaseous accretion.  The accretion process has been modeled with the inclusion of magnetic fields, electromagnetic radiation of the disk and energetic jets transporting angular momentum from the polar regions \\cite{Bardeen}-\\cite% {accretion}. Beside increasing the mass, accretion will also spin up the black holes. The spin limit reached due to accretion by canonical black holes (the system of a black hole and electromagnetically radiating accretion disk), as expressed in terms of the dimensionless spin is $\\chi _{can}=0.998$, close to the theoretically allowed maximum for a black hole, the unity. Indeed, such a high spin powering the jets seems compulsory for\\ understanding the low energy cutoff in the energetic electron spectra of jets in radio galaxies \\cite{cutoff}. Active Galactic Nuclei, in particular the closest, Cen A are the most likely sources of for the Ultra High Energy Cosmic Rays \\cite{UHECR}.  When galaxies merge, eventually their central SMBHs will also do so. Dynamical friction transfers some of the orbital angular momentum of the binary black hole system to the stellar environment, being ejected at the poles, a process which drives the system through the last parsec \\cite{Zier}% . Other mechanisms to overcome the last parsec are relaxation processes due to cloud/star -- star interactions, repopulating the stellar orbits in the center of the galaxy \\cite{Alexander}, binary orbital decay by three-body interactions in the gravitationally bound stellar cusps \\cite{Sesana}, or the interplay of three accretion disks: one around each black hole and the third, circumbinary, removing orbital angular momentum from the binary \\cite% {Hayasaki}.  At about $0.005$ parsecs gravitational radiation takes over dynamical friction as the leading dissipative effect \\cite{SpinFlip1}. For many SMBH binaries the gravitational waves emitted in the process of coalescence fall into the frequency range the long-delayed space mission LISA \\cite{LISA}. Depending on how rich in gas the binary environment may be and whether there any circumbinary disk has been formed, certain alignment between the proper spins and the orbital angular momentum could occur due to the Bardeen-Petterson effect (based in turn on the Lense-Thirring precession) \\cite{BaPe}. The two situations which could occur are the mergers precessing under random angle (also known as dry mergers) and non-precessing mergers, implying complete alignment of the spins and orbital angular momentum (wet mergers). The randomness in the orientation of precessing mergers typically reduces the final spin \\cite{HuBl}. For equal mass precessing mergers this varies from $0.69$ for non-spinning black holes up to values of $0.73$ for maximally spinning black holes \\cite{BertiVolonteri}. For a mass ratio $1:10$ the range of final spins opens up to the interval between $0.2$ and $0.83$, respectively, as function of the initial spins. For non-precessing mergers all configurations would practically conserve a high initial spin during the inspiral.  When the two black holes are at large distances, the orbital angular momentum $L$ is always much larger then the individual spins $S_{i}$. However at the characteristic radius $r^{\\ast }\\approx 0.005$ parsec% \\footnote{% The distance $r^{\\ast }$ depends weakly, as $m^{5/11}$ on the total mass and negligibly, as $\\eta ^{2/11}$ on the symmetric mass ratio $\\eta =\\mu /m=\\left( q^{1/2}+q^{-1/2}\\right) ^{-2}$, where $\\mu =m_{1}m_{2}/m$ is the reduced mass.} (and the corresponding post-Newtonian (PN) parameter $% \\varepsilon ^{\\ast }=Gm/c^{2}r^{\\ast }\\approx 10^{-3}$, defining the beginning of the inspiral, see \\cite{SpinFlip1}), where gravitational radiation takes over dynamical friction as the leading order dissipative effect, the ratios $S_{i}/L\\approx \\left( \\varepsilon ^{\\ast }\\right) ^{1/2}q^{3-2i}\\chi _{i}$ depend on the mass ratio $q\\geq 1$. For a maximally spinning larger black hole and separation $r_{\\ast }$ the ratio is one at about $q\\approx 30$. For mass ratios larger than $30$ therefore the spin dominates over the orbital angular momentum during the whole inspiral.  During the inspiral gravitational radiation further reduces the orbital angular momentum, but not the spin magnitudes. The spins will only precess driven by the leading order spin-orbit coupling and corrections due to spin-spin and mass quadrupole - mass monopole coupling \\cite{BOC}. In the process the direction of the total angular momentum remains unchanged, in an averaged sense over one radial orbit \\cite{ACST}.  Gravitational radiation does not modify this conclusion on short time-scales. Radiative evolutions with spin-orbit \\cite{GPV3}, spin-spin \\cite{spinspin1} and mass quadrupole - mass monopole \\cite{quadrup} couplings have been investigated, and their analysis in Ref. \\cite{spinspin2} lead to the important result that the instantaneous radiative changes of the spins average out during a radial period. On this timescale therefore there is no secular radiative change of the spin vectors at all:% \\begin{equation} \\left\\langle \\frac{d\\mathbf{S_{i}}}{dt}\\right\\rangle =0\\ .  \\label{nospinave} \\end{equation}% This result confirms that the spin dynamics can be regarded as a pure precession, up to high PN orders including radiation reaction, on the timescales comparable with a radial orbit. On much larger timescales however the spin vector will undergo a reorientation (spin-flip), as explained in detail in Ref. \\cite{SpinFlip1}.  There are several scenarios possible, according to the actual mass ratio::  a) The masses are comparable $m_{2}\\approx m_{1}$. In this case at the end of the inspiral (when the PN approximation breaks down) the orbital angular momentum still dominates over the spins \\cite{SpinFlip1}, \\cite{SpinFlip2}. Radiating away this remnant orbital angular momentum during the merger phase, while extrapolating the conservation of the direction of the total angular momentum and of the individual spin magnitudes to this phase \\cite% {finalspin}, \\cite{BR} could significantly reduce the final spin in all cases when the individual spins were severely misaligned with each other and with the orbital angular momentum. Such a misalignment would be typical in the case of precessing mergers. Therefore for equal masses in a precessing merger a not too high final spin can be considered typical.  b) The mass ratio is in the range $1:30$ to $1:3$. In this case the orbital angular momentum dominates over the spin only at the beginning of the inspiral, and as such is roughly aligned with the total angular momentum. At the end of the inspiral however the orbital angular momentum becomes smaller than the dominant spin, which has therefore to be reoriented towards the invariant total angular momentum direction. For precessing mergers this process causes a spin-flip during the inspiral, but does not reduce significantly the magnitude of the dominant spin \\cite{SpinFlip1}, \\cite% {SpinFlip2}. Non-precessing mergers on the other hand already imply an alignment of the spins and orbital angular momentum, therefore neither the spin magnitude, nor its direction will be changed by this process. In both cases, whatever happens to the orbital angular momentum during the plunge, its small value (compared to the dominant spin) at the end of the inspiral will obstruct any serious further change in the final spin.  In this mass ratio range therefore the magnitude of the dominant spin will not be much reduced by the merger, nevertheless a significant reorientation of its direction during the inspiral will typically occur for precessing mergers, which could be followed only by a minor further spin-flip during the plunge.  c) The mass ratio is less than $1:30$. Then the orbital angular momentum is too small from the beginning of the inspiral to modify the dominant spin. Neither its magnitude, nor its direction are affected and we practically face the inspiral of a test mass into the much larger black hole.  In this paper we revisit the merger process, based on the recent data of Ref. \\cite{CarameteBiermann}. We first derive the SMBH mass distribution. In Section \\ref{MassRatio} we fit a broken power law for the differential mass function, then, based on this fit and a number of simple and reasonable assumptions we derive the mass ratio distribution. We note that the results of Ref. \\cite{CarameteBiermann} are fully consistent with earlier results based on much smaller statistics \\cite{Greene}.  Next we derive in Section \\ref{FinalSpin} a simple approximant for the final spin of the emerging SMBH, as function of the mass ratio, initial spin magnitudes, and orientation of the spins with respect to the orbital plane. In the Appendix we compare the approximant with the more cumbersome expressions existing in the literature, which were obtained by fit to numerical simulations.  In Subsection \\ref{Dry} we adopt the configuration of precessing mergers, which allow for all relative spin and orbital angular momentum orientations on equal footing, lowering the chances for a large final spin after the merger. By integrating over all orientations in the precessing merger limit (without allowing any preference for alignment), for any initial spin set we determine a lower approximant to the final spin distribution as function of the mass ratio alone. By folding this with the previously derived mass ratio dependent merger rate, we obtain a \\textit{lower bound to the typical final spin} after SMBH mergers.  By contrast, in the non-precessing merger limit there is a perfect alignment of the spins with the orbital angular momentum, hence the integration should be carried on for this configuration alone, and only over the mass ratios, folded with the mass ratio distribution. By this method, in Subsection \\ref% {Wet} we get an \\textit{upper bound for the typical final spin}.  We discuss the implications of our results and present the concluding remarks in Section \\ref{Conclusions}. ",
        "conclusions": "}  In this paper we have studied the typical mass ratio and the typical final spin in a two-body system composed of supermassive black holes (SMBH), thus we did not consider the perturbations induced by either of the accretion disks, nearby stellar population, magnetic fields or jets. SMBHs reside in the center of each galaxy and following the frequent galaxy mergers they also merge. Various dissipative processes, like dynamical friction, accretion and emitted gravitational radiation are responsible to their gradual approach and finally gravitational radiation is which drives them to coalescence through a sequence of inspiral, merger and ringdown.  By starting from precise and new data on the SMBH mass distribution we derived both a differential and an integral mass function, shown on Fig. \\ref% {Fig_BH}. The differential mass function is a broken power law, with coefficients $\\,-1$ and $-3$, with the breakpoint at $8.9\\times 10^{7}$ M$% _{\\odot }$.  Then, exploiting a number of simple and reasonable assumptions we derived the mass ratio dependent probability of encounters of two such SMBHs, represented on Fig. \\ref{Fig_ratio}. This confirms our expectation that the most frequent (approximately half of the) encounters are for mass ratios $% 1:3 $ to $1:30$, the interesting mass ratio range where a spin-flip would occur during the inspiral \\cite{SpinFlip1}. \\begin{figure}[tbph] \\begin{center} \\centering\\includegraphics[width=8cm]{Fig4.eps} \\end{center} \\caption{(Color online) The typical final spin $\\protect\\chi _{f}$ as function of $\\protect\\chi _{1}=\\protect\\chi _{2}$ only, in the randomly precessing and the non-precessing merger limits (lower and upper curves, respectively). The curves are obtained by integration over all mass ratios of the expressions (\\protect\\ref{spinfindry}) $\\protect\\chi % _{f}^{prec}\\left( \\protect\\chi _{1}=\\protect\\chi _{2},q\\right) $ and (% \\protect\\ref{spinfinwet}) $\\protect\\chi _{f}^{align}\\left( \\protect\\chi _{1}=% \\protect\\chi _{2},q\\right) $, respectively, weighted with the mass ratio dependent probabilities of encounter (\\protect\\ref{N}). The line of equal initial and final spins is also indicated. Where this line crosses the final spin curves, there are two attractors (denoted by large dots), to where the final spin would converge after a sequence of mergers in the two scenarios.} \\label{Fig4} \\end{figure}  Next, based on certain well-founded assumptions we derived a simple analytical expression for the final spin of such a merger, depending on the mass ratio, initial spin magnitudes, and orientation of the spins with respect to the orbital plane and each other. This formula approximates well more cumbersome expressions obtained from the fit with numerical simulations, and it is much simpler, thus advantageous in order to carry on the cumbersome numerical integrations which followed.  We proceeded to find the typical final spin in two limiting and highly idealized scenarios. First we allowed for perfectly random orientations (precessing case), over which we have integrated, obtaining a final spin still depending on the initial spin magnitudes and mass ratio. Then we folded with the derived mass ratio dependent merger rate, we integrated over the mass ratio, deriving a lower bound to the typical final spin value after mergers.  Secondly we considered the non-precessing configuration, with all spins and the orbital angular momentum perfectly aligned. Folding the final spin for this particular configuration again with the derived mass ratio dependent merger rate and integrating over the mass ratio we obtained an upper bound for the typical spin. These are represented as function of the initial spin magnitudes (chosen to be equal\\footnote{% The second dimensionless spin parameter $\\chi _{2}$ will anyhow have but a small impact on the result, as the ratio of the \\textit{total} spins $S_{i}$ scales with the mass ratio squared.}) on Fig. \\ref{Fig4}. A third curve, the line of equal initial and final spins is also indicated on the figure. The fact that both slopes of $\\chi _{f}$ as function of the initial spins are smaller than one, leads to important consequences.  If we imagine a sequence of idealized (either randomly precessing or non-precessing) mergers, what happens is that low spins tend to increase by mergers while high spins decrease. There are in fact two attractors at $\\chi _{f}^{prec}=0.2$ and $\\chi _{f}^{align}=0.45$, respectively, where the spins converge after a reasonable number of the two types of mergers.  Real mergers, biased toward partial alignment by interactions with the environment (accretion, host galaxy, etc.) would generate a typical final spin lying between these two limiting values. Indeed, for example the galaxy group distribution around NGC383 \\footnote{This is also known as the 3C31 radio source, cf. the NASA Extragalactic Database.} looks like a spindle, with the spin of the central black hole in the dominant galaxy, NGC383, aligned with the long axis of the spindle. This shows a correlation between the dominant galaxy black hole spin and the distribution of the other galaxies, all with central black holes as well. It is to be expected that the distribution of galaxies in the environment of a dominant galaxy is not random, but correlated, such that in a merger of a galaxy with the dominant galaxy the final spin, depending on the nature of the correlation, could fall anywhere between the two curves shown of Fig. \\ref{Fig4}.  After the merger episode gaseous accretion can start to increase the spin again. If gaseous accretion were strong, then the spin could become quite large in relatively short time. We propose to work out quantitatively such a model in a forthcoming work."
    },
    "1208/1208.1644_arXiv.txt": {
        "abstract": "While the {\\it Helioseismic and Magnetic Imager} (HMI) onboard the {\\it Solar Dynamics Observatory} (SDO) provides Doppler velocity [$V$], continuum intensity [$I_C$], and line-depth [$Ld$] observations, each of which is sensitive to the five-minute acoustic spectrum, the {\\it Atmospheric Imaging Array} (AIA) also observes at wavelengths -- specifically the 1600 and 1700 Angstrom bands -- that are partly formed in the upper photosphere and have good sensitivity to acoustic modes. In this article we consider the characteristics of the spatio--temporal Fourier spectra in AIA and HMI observables for a 15-degree region around NOAA Active Region 11072. We map the spatio--temporal-power distribution for the different observables and the HMI Line Core [$I_L$], or Continuum minus Line Depth, and the phase and coherence functions for selected observable pairs, as a function of position and frequency. Five-minute oscillation power in all observables is suppressed in the sunspot and also in plage areas. Above the acoustic cut-off frequency, the behaviour is more complicated: power in HMI $I_C$ is still suppressed in the presence of surface magnetic fields, while power in HMI $I_L$ and the AIA bands is suppressed in areas of surface field but enhanced in an extended area around the active region, and power in HMI~$V$ is enhanced in a narrow zone around strong-field concentrations and suppressed in a wider surrounding area. The relative phase of the observables, and their cross-coherence functions, are also altered around the active region. These effects may help us to understand the interaction of waves and magnetic fields in the different layers of the photosphere, and will need to be taken into account in multi-wavelength local helioseismic analysis of active regions. ",
        "introduction": "\\label{intro}  Helioseismology relies on the observation of waves, as they affect light from the outer layers of the Sun, to infer the structure and dynamics of the otherwise invisible deeper layers. By studying the oscillations in multiple wavelengths of light -- and hence at different heights in the atmosphere -- simultaneously, it is possible to probe the layers of the atmosphere. Such investigations, aiming both to better understand the behaviour of the waves and their interaction with magnetic fields, and to improve the inference of subsurface properties, have a long history.  The launch in February 2010 of the {\\it Solar Dynamics Observatory} (SDO), carrying both the {\\it Helioseismic and Magnetic Imager} (HMI) and the {\\it Atmospheric Imaging Assembly} (AIA), provides new opportunities for cross-spectral helioseismic analysis, with full-disc, high-cadence images at many UV and EUV wavelengths from AIA as well as photospheric Doppler-velocity, continuum, and magnetic data from HMI. As discussed by \\citet{2011JPhCS.271a2058H}, the AIA {1600\\ \\AA} and {1700\\ \\AA} near-ultraviolet bands show a clear signature of the five-minute acoustic spectrum that  -- at least in the Sun-as-a-star case -- is much less contaminated by granulation ``noise'' than the continuum intensity in HMI's visible 6173~{\\AA} line.    Both the effect of surface magnetic activity on helioseismic waves in local areas and the phase and coherence relationships between the velocity and the intensity of radiation at different wavelengths have been studied for nearly two decades.    \\begin{figure} \\includegraphics[width=11.0cm]{fig1.ps} \\caption{ Spatial variation of each observable averaged over 23 May 2010: HMI~$V$ (a), HMI $I_C$ (b), AIA {1600 \\AA} (c), AIA {1700 \\AA} (d), HMI $I_L$ (e), and HMI longitudinal magnetic-field strength (f).} \\label{fig:1} \\end{figure}  \\begin{figure} \\includegraphics[width=11.0cm]{fig2.ps} \\caption{The mean magnetic inclination (a), from HMI vector magnetograms, courtesy K. Hayashi, an AIA {304\\ \\AA} image (b), and an AIA {171\\ \\AA} image (c), for the area of interest on 23 May 2010. The contours overlaid on panels a and b show the {500\\ G} level of the mean magnetic field strength from HMI line-of-sight magnetograms.} \\label{fig:1a} \\end{figure}   \\citet{1992ApJ...394L..65B} reported finding small areas associated with active regions that produced a disproportionate amount of acoustic power in the 5.5\\,--\\,7.5 mHz frequency band in ground-based Doppler observations using the Fe~{\\sc i}~{5576\\ \\AA} line, while \\citet{1992ApJ...392..739B} made the first observations of defined acoustic haloes around active regions in the same frequency band using Ca~K intensity. However, atmospheric seeing can have confusing effects on ground-based intensity observations \\citep{2001ESASP.464..219H}, particularly in sunspot umbrae. The picture became clearer with the launch of the {\\it Michelson Doppler Imager} \\citep[MDI:][]{1995SoPh..162..129S} onboard the {\\it Solar and Heliospheric Observatory} (SOHO), which allowed full-disc observations from space in both intensity and velocity using the {6768\\ \\AA} Ni line. \\citet{2002A&A...387.1092J} studied the power distribution in the line-depth, Doppler-velocity, and continuum-intensity observables from MDI. In both sunspot and plage regions they found that the power in areas of strong magnetic field  was suppressed in all observables, while for velocity and line depth (but not for continuum intensity) there was a ``halo'' of enhanced power surrounding the magnetic-field concentrations, which they concluded was acoustic in origin. The {\\it Transition Region and Coronal Explorer} (TRACE) made possible high-resolution, space-based observations in the 1600 and {1700\\ \\AA} bands. \\citet{2001ApJ...554..424J} used data from another SOHO instrument, {\\it Solar Ultraviolet Measurements of Emitted Radiation} (SUMER), TRACE, and MDI to investigate the oscillations in the chromosphere. This work established that the chromospheric modes were primarily the same $p$-modes seen in the photosphere. \\citet{2001A&A...379.1052K} used TRACE to study oscillations in the {1600\\ \\AA} UV band in the quiet Sun and found enhanced power at three-minute periods around patches of network field. \\citet{2003A&A...401..685M} also looked at acoustic power around an active region in short (two\\,--\\,four hour) stretches of TRACE data and did not find a high-frequency acoustic halo in the 1600 or {1700\\ \\AA} bands.  The {\\it Hinode/Solar Optical Telescope} has also been used for multi-wavelength seismology using the Ca~H line and the G band, for example by \\citet{2007PASJ...59S.637S}, who considered power, phase, and coherence spectra but not spatial maps, and by \\citet{2007PASJ...59S.631N} who looked at the power distribution in and around a sunspot. In the latter article there was no high-frequency acoustic halo detected, but there was a narrow band of excess power surrounding the sunspot at all wavelengths in the G band, and pronounced suppression of power at both wavelengths in the penumbra and plage, while the umbra showed excess power in the H line, especially above {4.5\\ mHz}.  \\citet{2011SoPh..268..349S} have examined the power distribution for velocity observations from MDI and found that the excess high-frequency power corresponds to regions with the magnetic-field inclination (as deduced from potential-field source-surface extrapolation) in the 40\\,--\\,60 degree range. However, the exact origin of these high-frequency haloes remains unclear.   The relative phase of the Doppler velocity and the intensity in various wavelength bands has been studied for more than two decades in both resolved and unresolved observations. Early (ground-based) spatially resolved phase measurements were made, for example, by \\citet{1990A&A...236..509D, 1992A&A...266..560D}; in the latter article they used three different wavelengths including the Na D line, formed in the lower chromosphere. These measurements revealed a complicated pattern of phase relationships that was not easily explained by simple models. \\citet{2007A&A...471..961M} made ground-based observations from the South Pole using the MOTH instrument \\citep{2004SoPh..220..317F}, taking high-cadence images in the sodium and potassium D lines. The sodium line shows clear evidence of a high-frequency halo around an active region.     \\begin{figure} \\includegraphics[width=11.0cm]{fig3.ps} \\caption{Slices through the smoothed relative-power maps for HMI $I_C$, at 3\\ mHz (a), 5\\ mHz (b), 7\\ mHz (c), 9\\ mHz (d), and in longitude and temporal frequency along the horizontal line shown in panels a\\,--\\,d (e). The grey-scale bar applies to all of the panels.} \\label{fig:2} \\end{figure} \\begin{figure} \\includegraphics[width=11.0cm]{fig4.ps} \\caption{Slices through the smoothed relative-power maps for HMI~$V$ at 3\\, mHz (a), 5\\ mHz (b), 7\\ mHz (c), 9\\ mHz (d), and in longitude and temporal frequency along the horizontal line shown in panels a\\,--\\,d (e). The grey-scale bar applies to all of the panels.}  \\label{fig:3} \\end{figure}  \\begin{figure} \\includegraphics[width=11.0cm]{fig5.ps} \\caption{Slices through the smoothed relative-power maps for HMI $I_L$, at 3\\ mHz (a), 5\\ mHz (b), 7\\ mHz (c), 9\\ mHz (d), and in longitude and temporal frequency along the horizontal line shown in panels a\\,--\\,d (e). The grey-scale bar applies to all of the panels.} \\label{fig:4} \\end{figure}  \\begin{figure} \\includegraphics[width=11.0cm]{fig6.ps} \\caption{Slices through the smoothed relative-power maps for {AIA 1700\\ \\AA}, at 3\\ mHz (a), 5\\ mHz (b), 7\\ mHz (c), 9\\ mHz (d), and in longitude and temporal frequency along the horizontal line shown in panels a\\,--\\,d (e). The grey-scale bar applies to all of the panels. } \\label{fig:5} \\end{figure}  \\begin{figure} \\includegraphics[width=11.0cm]{fig7.ps} \\caption{Slices through the smoothed relative-power maps for {AIA 1600\\ \\AA}, at 3\\ mHz (a), 5\\ mHz (b), 7\\ mHz (c), 9\\ mHz (d), and in longitude and temporal frequency along the horizontal line shown in panels a\\,--\\,d (e). The grey-scale bar applies to all of the panels.} \\label{fig:6} \\end{figure}       In this article we will examine the behaviour and relationships of the different observables in a small region of the solar surface containing an active region as observed by HMI and AIA, and show how the power and phase of the oscillations is affected by the presence of local magnetic fields. For this purpose we will show observations of NOAA Active Region 11072 on 23 May 2010. We have also examined the data for the same region over the two preceding and following days; the results are very similar and for reasons of space are not shown in this article.  In Section~\\ref{sec:2} we will describe our data and analysis and define the phase and coherence spectra. In Section~\\ref{sec:3} we show the region in each of the observables; in Section~\\ref{sec:4} we present the power maps for each observable, and in Section~\\ref{sec:5} we show the phase and coherence spectra. In Section~\\ref{sec:6} we discuss our findings. ",
        "conclusions": "\\label{sec:6} We have examined the spatio--temporal power distribution around an active region in a number of HMI and AIA observables, and the phase and coherence relationships between the intensity observables and the HMI Doppler velocity.  Five-minute power in all observables is suppressed in the sunspot (which is dark at all wavelengths) and also in plage areas (bright in AIA bands and HMI $I_L$). Above the acoustic cut-off frequency the behaviour is more complicated. Power in HMI $I_C$ is suppressed in the presence of surface fields at all frequencies, while power in HMI~$V$ is enhanced in a narrow zone around field concentrations (especially plage) and suppressed in a wider surrounding area; these results are consistent with earlier work using MDI. Power in HMI $I_L$ and the  AIA bands is suppressed in areas of surface field but -- in contrast to the results of \\citet{2003A&A...401..685M} -- enhanced in an extended area around the active region. In the HMI $I_L$ case, the pattern of narrow halo, encroaching on the region where the five-minute power is suppressed, and surrounding region of suppressed power, seen for the HMI~$V$ at {7\\ mHz}, is seen instead at {9\\ mHz}. For the UV bands, however, this pattern does not appear at any frequency; instead, the halo power fades with increasing frequency and vanishes above about {10\\ mHz} while the power suppression remains.  In all cases but HMI $I_C$, the regions of enhancement and suppression of power appear to move inwards towards the active region at increasing frequency.  The relative phase of the observables is altered around active regions. While the apparently in-phase behaviour of most observables in strong-field areas is associated with low coherence and thus of low significance, there are exceptions to this. In particular, in the 7~mHz frequency band there are areas close to the active region where the AIA {1700 \\AA} and {1600 \\AA} bands are both close in phase to the HMI velocity and coherent with it, whereas in the quiet Sun well away from the active region the phase difference is close to 90 degrees and the coherence is lower. This effect is not seen in the HMI $I_L$, which is otherwise very similar in its behaviour to the UV bands, although there is a phase difference between the two observables. At {9\\ mHz}, the 1600 and {1700\\ \\AA} bands also show qualitatively different behaviour, suggesting that though there may be substantial overlap between the two wavelength bands, the highest-frequency waves are sensing different layers of the atmosphere.  Clearly, within the ``halo'' zone surrounding the active region the propagation and reflection of the high-frequency waves differs from that over quiet Sun, whether due to alteration in the height of formation of the observables, thermal changes, or the direct effect of magnetic fields in trapping, scattering or transforming the waves; also, there are qualitative differences between the purely photospheric observables -- even at the relatively high formation level of the HMI $I_L$ -- and the UV bands.  In general, the higher the observable is formed, the higher the frequency at which the halo of excess power is found, and the more extensive the halo at the lowest frequency at which it appears. It is tempting to interpret these halo effects in terms of spreading magnetic-field canopies, as has been suggested, for example, by \\citet{2007A&A...471..961M}, but we note that the extent of the high-frequency halo does not exactly match the morphology of the overlying field seen in the EUV images. The size of the halo and the inner and outer regions of suppressed power appear to vary with frequency as well as with the height of formation of the observable used, contracting with increasing frequency. If the patterns of power distribution are related to spreading out of magnetic fields with height, this would imply that the higher-frequency pseudomodes are being reflected or absorbed at lower layers than the lower-frequency ones. The heights of formation themselves may also be altered in the presence of magnetic fields, further complicating the picture.  These effects no doubt also influence the changes in the local helioseismic parameters of modes above the acoustic cut-off frequency, seen for example by \\citet{2004ApJ...608..562H,% 2008ASPC..383..305H}, which tend to be in the opposite sense from those experienced by modes in the five-minute range.   Eventually, we hope to carry out multi-wavelength helioseismic analysis using HMI and AIA data together to extract parameters for local and global acoustic spectra. As the results described here demonstrate, any such analysis will need to take into account the effects of active regions on the phase and coherence of the oscillations.  The continuous, high-cadence, full-disc observations of HMI and AIA allow us to study the behaviour of waves in the photosphere and lower chromosphere at a level of detail that has not previously been possible. These observations may help us to improve our understanding of the interaction of waves and magnetic fields in the different layers of the photosphere.  While the analysis described here does not distinguish between true waves and other short-period spatio--temporal variations, we hope in the future to examine the three-dimensional power spectrum and its phase and coherence in order to separate the waves from granulation and other background effects.     \\begin{acks} RH thanks the National Solar Observatory for computing support, and Yvonne Elsworth for useful discussions. We also thank A.G. Kosovichev for useful comments on the manuscript. SDO data courtesy SDO (NASA) and the AIA and HMI consortia. This work was partly supported by NASA grant NNH12AT11I to NSO. This research has made use of NASA's Astrophysics Data System. \\end{acks}"
    },
    "1208/1208.1191_arXiv.txt": {
        "abstract": "We present a specific prescription for the calculation of cosmological power spectra, exploited here at two-loop order in perturbation theory (PT), based on the multi-point propagator expansion. In this approach power spectra are constructed from the regularized expressions of the propagators that reproduce both the resummed behavior in the high-$k$ limit and the standard  PT results at low-$k$. With the help of $N$-body simulations, we show that such a construction gives robust and accurate predictions for both the density power spectrum and the correlation function  at percent-level  in the weakly non-linear regime. We then present an algorithm that allows accelerated evaluations of all the required diagrams by reducing the computational tasks to one-dimensional integrals. This is achieved by means of pre-computed kernel sets defined for appropriately chosen fiducial models. The computational time for two-loop results is then reduced from a few minutes, with the direct method, to a few seconds with the fast one. The robustness and applicability of this method are tested against the power spectrum {\\tt cosmic emulator} from which a wide variety of cosmological models can be explored. The fortran program with which direct and fast calculations of power spectra can be done, \\RegPT, is publicly released as part of this paper. ",
        "introduction": "\\label{sec:intro} Since recombination, the large-scale structure of the Universe has evolved dominantly under the influence of both the cosmic expansion and the force of gravity acting on a pressure-less fluid. The statistical nature of its spatial clustering is therefore expected to bring valuable cosmological information about the dynamics of the cosmic expansion and structure formation. Of particular importance is the measurement of baryon acoustic oscillations (BAOs) imprinted on the power spectrum or two-point correlation function (e.g., \\cite{Eisenstein:2005su,Percival:2009xn,Blake:2011wn, Seo:2012xy,Anderson:2012sa}) from which one can precisely determine the cosmological distance to the high-redshift universe, and henceforth clarify the nature of late-time cosmic acceleration (e.g., \\cite{Seo:2003pu,Blake:2003rh,Glazebrook:2005mb,Shoji:2008xn, Padmanabhan:2008ag}). Precious  information regarding the growth of structure are and will also be obtained from redshift-space distortions (e.g., \\cite{Linder:2007nu,Guzzo:2008ac,Yamamoto:2008gr, Song:2008qt,Blake:2011rj}) and weak lensing measurements (see \\cite{2000A&A...358...30V,2008A&A...479....9F} and review papers \\cite{2001PhR...340..291B,2003astro.ph..5089V}) at scales ranging to the linear or quasi-linear to the non-linear regimes. This could be captured with unprecedented details  with the ongoing and future surveys, thanks to their redshift depth and large angular area, such as the Sloan Digital Sky Survey III\\footnote{{\\tt www.sdss3.org}}, the WiggleZ survey\\footnote{{\\tt wigglez.swin.edu.au}}, the Subaru Measurement of Imaging and Redsfhits\\footnote{{\\tt sumire.ipmu.jp/en/}}, the Dark Energy Survey\\footnote{{\\tt www.darkenergysurvey.org}}, the BigBOSS project\\footnote{{\\tt bigboss.lbl.gov/index.html}}, the Physics of the Accelerating Universe collaboration\\footnote{{\\tt www.pausurvey.org}} and the ESA/Euclid survey \\footnote{{\\tt www.euclid-ec.org}}.  With the  advent of such wealth of observations,  there is therefore a growing interest in the development of theoretical tools to accurately compute  the statistical quantities of the large-scale structure. At decreasing redshift and scale, the evolution of the large-scale structure however deviates significantly from the linear theory prediction and  non-linear gravitational clustering effects have to be taken into account. While $N$-body simulations can be relied upon in specific cases, because of the range of scales to be covered and the variety of models to explore, they should be complemented by analytical investigations that aim at computing the statistical properties of the large-scale structure from first principles, henceforth extending the validity range of linear calculations. It is to be noted that even at the scale of BAOs, linear calculations and one-loop standard PT corrections perform poorly (see e.g., \\cite{Crocce:2005xy,Carlson:2009it,Taruya:2009ir}) asking for more advanced PT calculations. The improvement of perturbation theory is thus a critical issue for the scientific exploitation of the coming surveys. Various  resummation schemes have been proposed in Refs. \\cite{Crocce:2005xy,Crocce:2005xz,Crocce:2007dt,Matsubara:2007wj, Matsubara:2008wx,McDonald:2006hf,Izumi:2007su,Taruya:2007xy, Taruya:2009ir, Pietroni:2008jx,Matarrese:2007wc,Valageas:2003gm, Valageas:2006bi} that aim at improving upon standard schemes. The aim of this paper is not to compare them but to propose, and test, a specific scheme that can be used routinely in practice.  In this paper, we are particularly interested in one of the resummation treatments, advocated in Ref.~\\cite{Bernardeau:2008fa}. In this approach, the standard PT expansion is re-organized by introducing the multi-point propagators. These are the ensemble average of the infinitesimal variation of the cosmic fields with respect to the initial conditions. A key property shown in the previous reference is that all the statistical quantities such as power spectra and bispectra can be re-constructed by an expansion series written solely in terms of the multi-point propagators. This is referred to as the multi-point propagator expansion or $\\Gamma$-expansion. The advantage of this approach is that the non-perturbative properties, which can be obtained in standard PT by summing up infinite series of PT expansions, are whole encapsulated in the multi-point propagators, including the effect of vertex renormalization. Furthermore, the $\\Gamma$-expansion has been found to be valid not only for Gaussian initial conditions, but also  for non-Gaussian ones~\\cite{Bernardeau:2010md}. The construction of accurate calculation scheme for power spectra and bispectra can then be split in pieces that can be tested separately.  The second key property that leads us to consider such objects is that their global shape, e.g. their whole $k$-dependence, can be computed in a perturbation theory context and compared to $N$-body results thanks to the high-$k$ exponential damping tail they all exhibit \\cite{Bernardeau:2008fa,Bernardeau:2011vy}. All these properties make the multi-point propagators the most important building blocks in the $\\Gamma$-expansion and the focus of our modeling efforts. In the following we will in particular make full use of the novel regularization scheme proposed in Ref.~\\cite{Bernardeau:2011dp} that allows to consistently  interpolate between  standard PT results at low-$k$ and the expected resummed behavior at high-$k$. This scheme has been explicitly tested for the two-point propagators up to two-loop order in  Ref.~\\cite{Bernardeauetal2012a} and for (specific shapes of) the three-point propagators in Ref.~\\cite{Bernardeau:2011dp}.   The first objective of this paper is to present an explicit calculation of the non-linear power spectrum and correlation function of the cosmic density field based on this regularized treatment. Of particular interest is the extent to which the proposed scheme for $\\Gamma$-expansion works beyond standard PT when corrections at next-to-next-to-leading, i.e. two-loop, order are included. Results will be checked with $N$-body simulations. We will see that the $\\Gamma$-expansion with the regularized treatment of propagators, which we hereafter call \\RegPT, has good convergence properties and agree remarkably well with simulations entirely covering the scales of BAOs at any redshift.   The  second objective of this paper is to design and exploit a method to accelerate the power spectrum computations. Power spectra calculations in the context of  \\RegPT~calculations are rather involved requiring multi-dimensional integrations that have to be done with time-consuming Monte Carlo calculations. Typically, computing the power spectrum at percent level from our scheme takes several minutes. While this is acceptable when a handful of models have to be computed, this is an obstacle when a large domain of parameter space has to be systematically explored. Making use of the $\\Gamma$-expansion functional form, we found though that it is possible to exploit a novel technique for accelerated calculation, in which only one-dimensional integrals need to be evaluated while ensuring the same precision as rigorous \\RegPT~calculations. The bottom line of this approach is to see the resulting nonlinear power spectrum as a functional of the linear power spectrum and then Taylor expand this form with respect to the linear spectrum shape.  We found that for well chosen fiducial models, it is sufficient to Taylor expand to first order only.  We are then led to prepare in advance a set of kernel functions encoding the \\RegPT~results for  well chosen fiducial models, whose normalizations are left floating, from which the \\RegPT~predictions for the target model can be calculated. We publicly release the fortran code, \\RegPT, as a part of this paper\\footnote{ The code is available at \\\\ {\\footnotesize{\\tt www-utap.phys.s.u-tokyo.ac.jp/\\~\\,ataruya/regpt\\_code.html}} }.   The organization of this paper is as follows. We begin by recalling the basic equations for cosmic fluid and perturbation theory in Sec.~\\ref{sec:RegPT}. We introduce the multi-point propagator and give the power spectrum expression based on the $\\Gamma$-expansion. With the regularized treatment of multi-point propagators, in Sec.~\\ref{sec:power_spec}, we examine the the power spectrum calculations including the corrections up to the two-loop order, and investigate their UV and IR sensitivity in evaluating the PT kernels. Then, in Sec.~\\ref{sec:comparison}, a detailed comparison between PT calculation and $N$-body simulation is presented, and the accuracy and range of validity of PT calculation is checked. Based on this, Sec.~\\ref{sec:PTreconst} describes in detail the method to accelerate the power spectrum calculations. Robustness and applicability of the accelerated \\RegPT~calculations to a wide range of cosmological models are tested against power spectrum {\\tt cosmic emulator} code in Sec.~\\ref{subsec:reliable}.   Finally, in Sec.~\\ref{sec:conclusion}, we conclude and explore practical extensions of this work. The description of the publicly released code, \\RegPT, is presented in Appendix~\\ref{appendix:RegPT_code}. ",
        "conclusions": "\\label{sec:conclusion}  It is needless to say that future cosmological observations make the development of cosmological tool aiming at accurately predicting the large-scale statistical properties of the universe highly desirable. In the first part  of the present paper, based on a renormalized perturbation theory (PT), we introduced an explicit computation scheme applied to the matter power spectrum and correlation function in weakly non-linear regime that consistently includes the PT corrections up to the two-loop order. The construction of the full expression for the power spectrum is based on the $\\Gamma$-expansion, i.e.  makes use of the multi-point propagators which are properly regularized so as to recover their expected resummed behavior at high-$k$ and to match the standard PT result at low-$k$. We call this regularized PT treatment \\RegPT. We have shown that the \\RegPT~scheme provides an accurate prediction for both the power spectrum and the correlation function, leading to a percent-level agreement with $N$-body simulations in the weakly non-linear regime.   In the second half of the paper, we presented a method to accelerate the power spectrum calculations. The method utilizes prepared data sets for some specific fiducial models from which regularized PT calculations can be performed for arbitrary cosmological models. The main interest of this method is that the evaluation the residual PT corrections between fiducial and target cosmological models can be reduced to mere one-dimensional  integrals. This enables us to dramatically  reduce the computational cost, and even with single-node calculation by a laptop computer, the power spectrum calculation can be done in a few seconds. We call this method \\RegPTfast, and we have demonstrated that the \\RegPTfast~treatment can perfectly reproduce the direct \\RegPT~calculations that involve several multi-dimensional integrals.  We then investigated the range of applicability of the \\RegPT~schemes in a broad class of cosmological models. For this purpose, we select 38 cosmological models, and compared the \\RegPT~predictions -- eventually incorporating the accelerated computations -- with results of a power spectrum emulator code, {\\tt cosmic emulator}. We show that with the help of  three fiducial models the \\RegPTfast~calculations give reliable predictions for the power spectra over this range of cosmological models\\footnote{Our analysis is however restricted to flat $w$-CDM models.}. We furthermore put forward an empirical criterion (\\ref{eq:k_crit}) that gives a good indication of the applicable range of the \\RegPT~scheme in $k$. The \\RegPTfast~treatment, together with the direct \\RegPT~calculation, has been implemented in a fortran code that we publicly release as part of this paper.   Although this paper is focused on precision calculations of the matter power spectrum, the \\RegPT~framework as well as the methodology for accelerated calculation can naturally be applied to the power spectrum of the velocity divergence and the cross-power spectrum of velocity and density fields in a similar way. The analysis of the velocity power spectrum, together with a detailed comparison with $N$-body simulations, will be presented elsewhere. Of particular interest is the application of the \\RegPT~schemes  to the redshift-space power spectrum or correlation function. In this case, not only the velocity and density power spectra, but also the multi-point spectra like bispectrum, arising from the non-linear mode coupling, seem to play important roles, and should be properly modeled. Significance of the effect of multi-point spectra has been recently advocated by Refs.~\\cite{Reid:2011ar,Taruya:2010mx,Nishimichi:2011jm,Tang:2011qj}, and there appear physical models that account for this. Combination of these models with the \\RegPT~schemes would be very important, and we will discuss it in a near future."
    },
    "1208/1208.3194_arXiv.txt": {
        "abstract": "We use images of high spatial, spectral and temporal resolution, obtained using both ground- and space-based instrumentation, to investigate the coupling between wave phenomena observed at numerous heights in the solar atmosphere. Analysis of $4170${\\,}{\\AA} continuum images reveals small-scale umbral intensity enhancements, with diameters $\\sim$$0{\\,}.{\\!\\!}{\\arcsec}6$, lasting in excess of $30$~minutes. Intensity oscillations of $\\approx$$3$~minutes are observed to encompass these photospheric structures, with power at least three orders-of-magnitude higher than the surrounding umbra. Simultaneous chromospheric velocity and intensity time series reveal an $87\\pm8\\degr$ out-of-phase behavior, implying the presence of standing modes created as a result of partial wave reflection at the transition region boundary. We find a maximum wave guide inclination angle of $\\approx$$40\\degr$ between photospheric and chromospheric heights, combined with a radial expansion factor of $<$$76$\\%. An average blue-shifted Doppler velocity of $\\approx$$1.5$~km{\\,}s$^{-1}$, in addition to a time lag between photospheric and chromospheric oscillatory phenomena, confirms the presence of upwardly-propagating slow-mode waves in the lower solar atmosphere. Propagating oscillations in EUV intensity are detected in simultaneous coronal fan structures, with a periodicity of $172\\pm17$~s and a propagation velocity of $45\\pm7$~km{\\,}s$^{-1}$. Numerical simulations reveal that the damping of the magneto-acoustic wave trains is dominated by thermal conduction. The coronal fans are seen to anchor into the photosphere in locations where large-amplitude umbral dot oscillations manifest. Derived kinetic temperature and emission measure time-series display prominent out-of-phase characteristics, and when combined with the previously established sub-sonic wave speeds, we conclude that the observed EUV waves are the coronal counterparts of the upwardly-propagating magneto-acoustic slow-modes detected in the lower solar atmosphere. Thus, for the first time, we reveal how the propagation of $3$~minute magneto-acoustic waves in solar coronal structures is a direct result of amplitude enhancements occurring in photospheric umbral dots. ",
        "introduction": "Early white-light eclipse photographs of the Sun revealed elongated, faint columns of enhanced density stretching far out into the corona \\citep{van50, Sai65}. These structures, now commonly referred to as coronal plumes, can be viewed over a wide range of wavelengths, in particular the extreme ultraviolet \\citep[EUV;][]{Boh75, Ahm77}. Plumes are just one type of EUV feature that are seen to exist in the solar corona. Other examples include coronal loop and fan structures, which are observed to outline the coronal magnetic-field topology, and demonstrate a wide range of oscillatory behaviour \\citep[{\\rmfamily e.g.,}][]{Asc99, Asc02, Nak99, Jes08c, Ofm08, Ver08, Van08, Bal11}.  One of the first studies which uncovered propagating wave phenomena in coronal structures was that by \\citet{Def98}. These authors utilised the Extreme-ultraviolet Imaging Telescope on board the Solar and Heliospheric Observatory spacecraft to identify quasi-periodic perturbations in the brightness of $171$~{\\AA} images. \\citet{DeM00} undertook a similar study using the higher spatial resolution Transition Region and Coronal Explorer \\citep[TRACE;][]{Han99}, and concluded that these oscillations were signatures of slow magneto-acoustic waves, which propagate upwards along the coronal waveguides with velocities of $70$--$165$~km{\\,}s$^{-1}$ and periods in the range $180$--$420$~s. Energy estimates for these motions exhibit an incredibly wide range of values, typically $10^{2}$--$10^{5}$~ergs{\\,}cm$^{-2}${\\,}s$^{-1}$ \\citep{Def98, DeM00}.  Since magnetic fields play an essential role in plume/fan/coronal-loop formation and structuring, they are often modelled using magnetohydrodynamic (MHD) equations \\citep{Del97}. Utilising non-linear, two-dimensional MHD simulations, \\citet{Ofm99} were able to replicate previous observational results, and concluded that outward-propagating slow magneto-acoustic waves may be able to contribute significantly to the heating of the lower corona through compressive dissipation. Furthermore, theoretical modelling has suggested that the propagation characteristics of a magneto-acoustic intensity perturbation depends on a number of factors, including the dissipation of the wave energy \\citep{Kli04, DeM12}. Numerical simulations indicate that thermal conduction may be the dominant damping mechanism behind the dissipation of magneto-acoustic wave energy in the solar corona \\citep{Ofm02, DeM03, DeM04, Men04}.  \\begin{figure*} \\includegraphics[width=\\textwidth]{fig01_lowres.eps} \\caption{Simultaneous images of the lower solar atmosphere, obtained through $4170${\\,}{\\AA} continuum (left), H$\\alpha$ core (middle), and Ca~{\\sc{ii}}~$8542${\\,}{\\AA} core (right) filters at $13$:$32$~UT on $2011$ July $13$. The Ca~{\\sc{ii}}~$8542${\\,}{\\AA} image displays Doppler-compensated intensities, while the dashed white box outlines the ROSA/HARDcam field-of-view. Axes are in heliocentric arcseconds, where $1${\\arcsec} $\\approx$ $725$~km on the solar surface.} \\label{fig1} \\end{figure*}   Through examination of coronal structures in the close proximity of active regions, \\citet{Flu01} and \\citet{Mar06} were able to reveal how structures situated above sunspot regions displayed intensity oscillations with a period of the order of three minutes, while oscillations in ``non-sunspot'' structures demonstrated much longer periodicities \\citep{DeM02}. The authors concluded that the most likely explanation for the observed longitudinal waves revolves around a driver directly exciting the magnetic footpoints. This scenario requires the magneto-acoustic wave trains to be able to propagate from the lower solar atmosphere, through the transition region, and into the corona. Utilising numerical simulations of the Sun's lower atmosphere, \\citet{Kho06} have revealed how longitudinal oscillations, generated with a periodicity of $\\sim$$3$~minutes in sunspot umbrae, can readily propagate upwards from the photosphere and into the chromosphere.  Previously, three minute umbral oscillations have been notoriously difficult to detect at photospheric heights. \\citet{Bal87} were unable to detect any photospheric signatures using the Locarno solar station at the G{\\\"o}ttingen Observatory, while \\citet{Bru85} suggested that they may get lost in the noise, as a result of their very low amplitude. Indeed, \\citet{Nag07} utilised Hinode/SOT image sequences to show how oscillatory power, at all frequencies, is significantly reduced in sunspot umbrae. More recently, \\citet{Kob08, Kob11} have not only detected photospheric three minute oscillations, but the authors also claim that the location of maximum chromospheric power also corresponds to a co-spatial decrease in power of the photospheric oscillations. However, the spatial resolution obtained by \\citet{Kob08, Kob11} was on the order of 1{\\arcsec}, so precise diagnostics of the exact umbral structures displaying three minute periodicities was impossible.  In this paper, we utilise ground- and space-based instrumentation, with high spatial, temporal and spectral resolution, to investigate the origin of $3$~minute magneto-acoustic waves observed in EUV images of coronal fan structures. We employ a multi-wavelength approach to study the photospheric counterpart of these coronal phenomena, and analyse the resulting wave propagation characteristics from the photosphere, through the chromosphere, and out into the corona.  \\newpage ",
        "conclusions": "Here we present high-cadence observations of the solar atmosphere, obtained using the latest ground- and space-based facilities. Prominent oscillatory behaviour is detected throughout the optical and EUV image sequences, with remarkable similarities found between the detected wave modes. First, a number of UD structures in the photospheric umbra are found to display intensity oscillations with a $\\approx$$3$~minute periodicity. These oscillations exhibit considerable power, with regions encompassing the UDs displaying more than three orders-of-magnitude stronger power than the background umbra (Figure~\\ref{fig3}). Next, chromospheric intensity {\\it{and}} velocity measurements were analysed for the presence of co-spatial and co-temporal oscillations. Such phenomena were detected, both on larger spatial scales, and with small central offsets with respect to the underlying photospheric oscillations. By considering the extension of magnetic flux tubes from the solar surface out into the upper solar atmosphere, a geometric expansion of only $76$\\% in the radial direction and an inclination angle $<$$40\\degr$ allows the observed oscillations to be interpreted as originating from within the same magnetic flux tube.  Following examination of the phase lag between chromospheric velocity and intensity components, a $V-I$ phase angle of $-87\\pm8\\degr$ was derived, allowing these waves to be described as a magneto-acoustic mode, with characteristics consistent with standing acoustic modes (middle panel of Figure~\\ref{fig5}). The generation of a chromospheric standing wave may be the result of partial wave reflection at the transition region boundary. An average blueshift velocity of $\\approx$$1.5$~km{\\,}s$^{-1}$ was found in the locations where high chromospheric oscillatory power was present. This is a sub-sonic velocity, and coupled with a time lag between photospheric and chromospheric oscillatory phenomena, strengthens our interpretation that the observed oscillations are upwardly-propagating magneto-acoustic waves, which originate in UD structures located in the photospheric umbra.  A prominent fan structure is present in the simultaneous coronal EUV images, namely those from the AIA $131${\\,}{\\AA}, $171${\\,}{\\AA}, $193${\\,}{\\AA}, $211${\\,}{\\AA}, and $335${\\,}{\\AA} bandpasses. This fan is not readily apparent in either the transition-region dominated $304${\\,}{\\AA} emission, or in the higher temperature ($\\sim$$7.0$~MK) $94${\\,}{\\AA} bandpass. Using DEM techniques, we constrained the temperature of the coronal fan to $0.5$ -- $1.2$~MK, thus placing it outside the temperature range of both the $94${\\,}{\\AA} and $304${\\,}{\\AA} filtergrams (Figure~\\ref{fig6}). Time-distance techniques were employed on the EUV imaging data where the fan structure was readily apparent, allowing the characteristics of propagating wave phenomena to be uncovered. Most coronal channels, regardless of their absolute temperature sensitivity, revealed outwardly propagating wave fronts with an average periodicity and velocity of $172\\pm17$~s and $45\\pm7$~km{\\,}s$^{-1}$, respectively (lower panel of Figure~\\ref{fig7}). The out-of-phase nature between the derived temperature and emission measure signals indicates the presence of a compressive wave mode (Figure~\\ref{fig8}). This, coupled with a sub-sonic wave speed ($\\approx$$45$~km{\\,}s$^{-1}$), highlights the fact that these coronal phenomena are best described as upwardly-propagating magneto-acoustic slow mode waves.  Employing numerical simulations, we were able to accurately simulate the behaviour of the coronal EUV emission. Utilising input parameters derived directly from the AIA observations, forward-modelling techniques allowed us to evolve velocity, density, and emissivity values forward in time, creating a time series which could be directly compared to the AIA observations. Crucially, our simulations revealed that thermal conduction is the primary damping mechanism behind the dissipation of magneto-acoustic slow-mode waves in the corona. Other mechanisms, including optically thin radiation and compressive viscosity, play a secondary role in the damping of these oscillations.  The fan structure observed in the AIA images, which displays signatures of propagating magneto-acoustic waves, appears to have anchor points in the south-west quadrant of the photospheric sunspot umbra (lower panels of Figure~\\ref{fig2}). These locations are also consistent with the presence of large-amplitude wave phenomena detected in simultaneous photospheric and chromospheric image sequences. The co-temporal and co-spatial relationship between these upwardly-propagating magneto-acoustic wave modes, detected throughout the entire solar atmosphere, suggests such coronal phenomena may be driven by UD oscillations occurring inside the sunspot umbra. With this conclusion, it appears that photospheric structures which are on the order of $0{\\,}.{\\!\\!}{\\arcsec}5$ ($360$~km) in diameter, can have a strong influence on coronal structures not only several thousand km above their position, but on structures which have expanded into the local plasma to diameters often exceeding $10${\\arcsec} ($7000$~km)."
    },
    "1208/1208.1228_arXiv.txt": {
        "abstract": "We present two millisecond pulsar discoveries from the PALFA survey of the Galactic plane with the Arecibo telescope. PSR~J1955+2527 is an isolated pulsar with a period of 4.87~ms, and PSR~J1949+3106 has a period of 13.14~ms and is in a 1.9-day binary system with a massive companion. Their timing solutions, based on 4 years of timing measurements with the Arecibo, Green Bank, Nan\\c{c}ay and Jodrell Bank telescopes, allow precise determination of spin and astrometric parameters, including precise determinations of their proper motions. For PSR~J1949+3106, we can clearly detect the Shapiro delay. From this we measure the pulsar mass to be $1.47^{+0.43}_{-0.31}$~\\msun, the companion mass to be $0.85^{+0.14}_{-0.11}$~\\msun\\ and the orbital inclination to be $i = 79.9_{+1.6}^{-1.9}$ degrees, where uncertainties correspond to $\\pm 1$-$\\sigma$ confidence levels. With continued timing, we expect to also be able to detect the advance of periastron for the J1949+3106 system. This effect, combined with the Shapiro delay, will eventually provide very precise mass measurements for this system and a test of general relativity. ",
        "introduction": "In this paper, we discuss timing results for two millisecond pulsars (MSPs) discovered by the PALFA Consortium\\footnote{\\tt http://www.naic.edu/alfa/pulsar} with the Arecibo telescope. The PALFA survey of the Galactic plane and follow-up observations of new discoveries are motivated by the wide applications of pulsar timing in exploring the composition, internal structure, and magnetospheric state of neutron stars. Millisecond pulsars in particular tend to be extremely stable rotators, which can be used to address a variety of problems in fundamental physics and astrophysics. One such outstanding problem is detecting gravitational waves and studying the properties of gravitational wave radiation from various types of sources and various epochs in the Universe's lifetime. Timing observations of MSPs in binary systems can be used to estimate the pulsar and companion masses via measuring post-Keplerian binary parameters. Neutron star mass measurements allow us to constrain the equation of state (EoS) of matter at densities larger than that of an atomic nucleus.   Until recently, all precise measurements of neutron star masses fell within a narrow range around 1.4~\\msun, the Chandrasekhar limit. However, recent, precise measurements of the masses of some MSPs showed that they can have significantly higher masses. PSR~J1903+0327 \\citep{Champion08}, the first MSP found in the PALFA survey, has a mass of $1.67 \\pm 0.02~\\msun$ (99.7 \\% confidence limit, \\citealt{Freire11}) which is significantly above the Chandrasekhar limit. PSR J1614$-$2230 was found to have a mass of $1.97 \\pm 0.04~\\msun$ \\citep{Demorest10}. These high masses rule out many EoSs for matter at densities higher than that of the atomic nucleus. In particular, the mass measurement for J1614$-$2230 rules out or highly constrains most proposed hyperon or boson EoSs (\\citealt{Glendenning98}, \\citealt{Lackey06}, \\citealt{Schulze06}, \\citealt{Lattimer07}).   In this paper, we describe two MSPs discovered by the PALFA survey, which uses the Arecibo telescope and the seven-beam ALFA receiver \\citep{Cordes06}. In addition to presenting full timing solutions for the two pulsars, we explore what intrinsic or extrinsic effects account for the overall TOA uncertainty in both cases. One of the two discoveries reported here is in a nearly edge-on binary system with a massive companion, allowing measurement of the Shapiro delay and consequently estimation of the pulsar and companion masses. ",
        "conclusions": "To the current tally of almost 2000 known pulsars we have added two MSPs found by the PALFA survey, J1949+3106 and J1955+2527, and presented their timing solutions. While J1955+2527 is isolated, in the J1949+3106 binary system we have been able to confidently measure the Shapiro delay and estimate the pulsar and companion masses. The pulsar's current median mass is 1.47~\\msun, and the companion's median mass is 0.85~\\msun. The uncertainties of these mass estimates are still $0.3 - 0.4~\\msun$ and $0.1 - 0.2~\\msun$, respectively, but they will improve with continued timing. We are also on the verge of being able to detect the relativistic periastron advance in this system.  We have outlined the steps towards breaking down the various contributions to the overall rms timing residual and applied them to the TOA sets of the two new pulsars. This type of characterization is important with a view to figuring out what effects and in what cases are mostly responsible for the observed residuals, and what can we do to mitigate them. This has further implications for how we can take maximum advantage of the properties of new discoveries in projects like the International Pulsar Timing Array (IPTA), which need extremely precise timing measurements on pulsars that are very stable natural clocks. The North American Nanohertz Observatory for Gravitational Waves (NANOGrav\\footnote{\\texttt{http://nanograv.org}}), an IPTA member, uses the Arecibo telescope for timing pulsars suitable for the IPTA and is especially interested in new MSP discoveries in the Arecibo sky. In the case of J1949+3106 and J1955+2527, the rms timing residuals are too large for including these pulsars in the IPTA sample. Table~\\ref{tab_toaerror} lists overall rms timing residuals for both pulsars, residuals by observatory, the expected contribution to the residual from radiometer noise, and upper limits on the contributions from pulse jitter, diffractive scintillation, and unmodeled DM variations. We find that radiometer noise puts a hard limit on TOA precision in both cases. The overall timing residual of J1949+3106 is consistent with that limit, while the residual of J1955+2527 is more than twice as large. The TOA residuals of J1949+3106 are consistent with white noise, while J1955+2527 exhibits some modest departures from white noise as evidenced in a zero-crossing test and in a fit for a second frequency derivative."
    },
    "1208/1208.1612_arXiv.txt": {
        "abstract": "We estimate the relative contributions of the supermassive black hole (SMBH) accretion disk, corona, and obscuring torus to the bolometric luminosity of Seyfert galaxies, using \\spitzer\\ mid-infrared (MIR) observations of a complete sample of 68 nearby active galactic nuclei (AGNs) from the \\integral\\ all-sky hard X-ray (HX) survey. This is the first HX-selected (above 15~keV) sample of AGNs with complementary high angular resolution, high signal to noise, MIR data. Correcting for the host galaxy contribution, we find a correlation between HX and MIR luminosities: $\\Lmir\\propto\\Lint^{0.74\\pm0.06}$. Assuming that the observed MIR emission is radiation from an accretion disk reprocessed in a surrounding dusty torus that subtends a solid angle decreasing with increasing luminosity (as inferred from the declining fraction of obscured AGNs), the intrinsic disk luminosity, $\\Ld$, is approximately proportional to the luminosity of the corona in the 2--300~keV energy band, $\\Lc$, with the $\\Ld/\\Lc$ ratio varying by a factor of 2.1 around a mean value of 1.6. This ratio is a factor of ${\\sim} 2$ smaller than for typical quasars producing the cosmic X-ray background (CXB). Therefore, over three orders of magnitude in luminosity, HX radiation carries a large, and roughly comparable, fraction of the bolometric output of AGNs. We estimate the cumulative bolometric luminosity density of local AGNs at ${\\sim} (1-3)\\times 10^{40}$~erg~s$^{-1}$~Mpc$^{-3}$. Finally, the Compton temperature ranges between $kT_{\\rm c}\\approx 2$ and $\\approx 6$~keV for nearby AGNs, compared to $kT_{\\rm c}\\approx 2$~keV for typical quasars, confirming that radiative heating of interstellar gas can play an important role in regulating SMBH growth. ",
        "introduction": "\\label{s:intro}  Active galactic nuclei (AGNs) are extremely powerful sources of electromagnetic radiation over many decades in frequency from radiowaves to gamma-rays. According to the commonly accepted scenario, an AGN shines due to accretion of gas onto a supermassive black hole (SMBH) residing in a galactic nucleus.  In Seyfert galaxies and quasars, most of the luminosity is emitted in the form of ultraviolet (UV) radiation generated in a geometrically thin, optically thick accretion disk \\citep{shasun73}, giving rise to a ``big blue bump'' (BBB) in the spectral energy distribution (SED, e.g., \\citealt{malsar82}). Additional, higher energy radiation is generated in a hot corona of the accretion disk (e.g., \\citealt{haamar93}) and possibly also in collimated outflows (jets) of relativistic plasma, producing a hard X-ray (HX) peak in the SED. The integrated (and redshifted) HX emission of all AGNs in the observable Universe makes up the bulk of the cosmic X-ray background. There is also a third, mid-infrared (MIR) peak in AGN SEDs (e.g., \\citealt{barvainis87}), which arises from reprocessing of a significant fraction of the disk's and some of the coronal radiation in a torus of molecular gas and dust surrounding the inner accretion flow. In fact, only in unobscured or ``type 1'' AGNs can all three spectral components, the HX bump, the BBB, and the MIR bump, be observed.  According to the unified model \\citep{ant93}, these are objects viewed through the funnel of the dusty torus. In contrast, only the HX and MIR components are visible in the SEDs of obscured or ``type 2'' AGNs, because the torus is opaque to UV emission from the accretion disk but transparent to coronal radiation at energies above $\\sim$15~keV (except in Compton-thick sources) and to its own infrared emission (at least at wavelengths $\\gtrsim20$~$\\mu$m). All other emission components, including broad- and narrow-line emission and non-thermal radio and gamma-ray radiation are usually not significant as regards their contribution to the angular-integrated bolometric luminosity of AGNs; these components will therefore not be discussed below.  To understand how electromagentic radiation is emitted and reprocessed during accretion of matter onto SMBHs, it is crucial to explore i) in what proportion the AGN luminosity is shared between the accretion disk and its corona, ii) what fraction of the bolometric luminosity is reprocessed in the torus, and iii) how these properties depend on black hole mass and accretion rate. One also needs such information to study the role of AGN feedback in regulating SMBH growth and galactic evolution. One of the proposed feedback mechanisms is photoionization and Compton heating of interstellar gas by AGN radiation (e.g., \\citealt{cioost01,proetal08}), whose efficiency critically depends on the AGN SED \\citep{sazetal04,sazetal05}. Finally, information on AGN SEDs can be used to derive bolometric corrections required to reconstruct the cosmic history of SMBH accretion growth based on AGN statistics provided by extragalactic surveys (e.g., \\citealt{maretal04,merhei08}).  Among all types of AGNs, the SEDs of unobscured high-luminosity quasars have been studied most extensively (see, e.g., \\citealt{elvetal94,ricetal06,shaetal11}). Their obscured counterparts -- type~2 quasars -- have been explored to a much lesser degree, although recent surveys have begun to find such objects in significant numbers (e.g., \\citealt{poletal06,hicetal07,lanetal09}). There is also much uncertainty with respect to the SEDs of Seyfert galaxies, which are typically less luminous than more distant quasars. The difficulty is that even in Seyfert 1s, the accretion disk emission is usually contaminated by host galaxy stellar emission in visible bands and the BBB peaks in the observationally difficult far-UV band (see, however, \\citealt{scoetal04,vasfab07,vasfab09}).  The goal of the present study is to systematically assess the relative contributions of the accretion disk, hot corona, and obscuring torus to the bolometric luminosity of local Seyfert galaxies. To this end, we i) cross-correlate the HX luminosities of nearby AGNs detected during the all-sky survey of the {\\em International Gamma-Ray Laboratory} (\\integral, \\citealt{winetal03}) with the MIR luminosities of these objects measured by the {\\em Spitzer Space Telescope} \\citep{weretal04}, and ii) use the proportion of obscured to unobscured AGNs to estimate the opening angle of dusty tori as a function of luminosity. We then put our findings for nearby AGNs into the broader context of cosmic SMBH growth by making a comparison with distant quasars.  Most previous relevant studies were based on AGN samples compiled in a fairly arbitrary manner from optical and/or soft X-ray (below 10~keV) catalogs (e.g., \\citealt{lutetal04,horetal06,honetal10}). In these energy bands, AGNs can easily be missed due to absorption, as powerful sources can become invisible when obscured by large amounts of dust and cold gas in the torus and/or host galaxy. Furthermore, as already noted above, optical emission from relatively low-luminosity AGNs can be diluted against the background of a luminous galaxy (see \\citealt{mushotsky04} for a detailed discussion of AGN selection at different wavelengths).  The hard X-ray band, above $\\sim 15$~keV, provides a census of AGNs that is far less biased with respect to the viewing orientation of the torus and is unbiased with respect to host galaxy properties. There have been a few previous attempts \\citep{vasetal10,muletal11} of systematically studying the MIR properties of HX selected AGNs using the {\\em Swift} all-sky hard X-ray survey \\citep{tueetal10}. However, these studies either used data from the IRAS all-sky photometric infrared survey, so that it was impossible to reliably subtract the host galaxy contribution from the AGN emission, or used high angular resolution \\spitzer\\ data but only for statistically incomplete subsamples of {\\em Swift} AGNs. Our \\integral\\ sample is the first statistically complete, HX selected sample of AGNs with complementary high angular resolution, high signal to noise, MIR data. The extensive \\spitzer\\ coverage (3.6--38~$\\mu$m) available for the entire \\integral\\ sample makes this a unique data set for studying SEDs of AGNs in the local Universe. ",
        "conclusions": ""
    },
    "1208/1208.6204_arXiv.txt": {
        "abstract": "We present the first resolved map of plane-of-sky magnetic field strength for a quiescent molecular cloud. GRSMC 45.60+0.30 subtends 40~$\\times$~10 pc at a distance of 1.88~kpc, masses 16,000~$M_{\\sun}$, and exhibits no star formation. Near-infrared background starlight polarizations were obtained for the Galactic Plane Infrared Polarization Survey using the 1.8m Perkins telescope and the Mimir instrument. The cloud area of 0.78 square degrees contains 2,684 significant starlight polarizations for 2MASS-matched stars brighter than 12.5~mag in $H$-band. Polarizations are generally aligned with the cloud's major axis, showing an average P.A. dispersion of $15^{\\circ}\\pm2^{\\circ}$ and polarization of $1.8\\pm0.6\\%$. The polarizations were combined with Galactic Ring Survey $^{13}$CO spectroscopy and the Chandrasekhar-Fermi method to estimate plane-of-sky magnetic field strengths, with an angular resolution of 100 arcsec. The average plane-of-sky magnetic field strength across the cloud is $5.40\\pm0.04$~$\\mu$G. The magnetic field strength map exhibits seven enhancements, or `magnetic cores.' These cores show an average magnetic field strength of $8.3\\pm0.9$~$\\mu$G, radius of $1.2\\pm0.2$~pc, intercore spacing of $5.7\\pm0.9$~pc, and exclusively subcritical mass-to-flux ratios, implying their magnetic fields continue to suppress star formation. The magnetic field strength shows a power law dependence on gas volume density, with slope $0.75\\pm0.02$ for $n_{H_2}$~$\\geq10$~cm$^{-3}$. This power law index is identical to those in studies at higher densities, but disagrees with predictions for the densities probed here. ",
        "introduction": "What roles do magnetic fields play in molecular clouds? Within clouds, important forces include gravity, gas pressure, cosmic ray pressure, and magnetic fields; however, they are not independent of each other. For example, \\citet{HC05} point out that cosmic rays and gas pressure are coupled through the magnetic field. Therefore, the magnetic field may be a major factor in cloud dynamics across many scales and may be a key agent in regulating the rate of star formation \\citep{FIE65}. Yet, the magnetic field is difficult to sense and harder still to map in much detail.  In the hot ionized interstellar medium, magnetic fields are probed in diffuse regions through Faraday rotation of background pulsars \\citep{SM68,MAN74,HAN06} and extragalactic sources \\citep{COO62,PS11}, as well as through synchrotron emission \\citep{WES62,WIE62,YUS84,YUS89}. In cold, denser regions, the magnetic fields are probed through Zeeman splitting of spectral lines \\citep{VER68,CRUT3}, polarized thermal emission from aligned dust \\citep{HIL88,DOT10}, and polarization of background starlight by aligned dust \\citep{MF70,HEI00}. Zeeman splitting and Faraday rotation measure the line-of-sight magnetic field strength; the other methods trace the orientation of the magnetic field on the plane of the sky. Chandrasekhar \\& Fermi (1953a; hereafter CF) derived an estimate of the plane-of-sky magnetic field strength using the dispersion of the starlight polarization position angles (P.A.s), the local gas motions, and the local gas density.  Polarization of background starlight, first observed by \\citet{HILT} and \\citet{HALL}, has been attributed to elongated dust grains aligned by local magnetic fields. Early alignment mechanisms were proposed by \\citet{DG51}, \\citet{GOLD}, and \\citet{PS71}. \\citet{L03} provides a comprehensive history of grain alignment.  Currently, the favored model for dust alignment is the radiative torque mechanism \\citep{DG76,L03,L07}. With radiative torques, spinning anisotropic dust grains preferentially align their long axes perpendicular to the local magnetic field. Unpolarized background starlight sees a higher extinction cross-section perpendicular to the local magnetic field direction and the light picks up a small linear polarization parallel to the projected magnetic field direction. With this mechanism operating, magnetic fields can be probed within molecular clouds using optical or near-infrared polarimetric observations of background stars.  The Galactic Plane Infrared Polarization Survey \\citep[GPIPS;][]{DAN11}, from which these data are drawn, is a near-infrared, linear polarization survey that measures the polarization of background starlight in the inner Galactic mid-plane ($18\\degr < \\ell < 56\\degr$, $-1\\degr < b < 1\\degr$), to and beyond the nearest spiral arm. This survey obtains $H$-band (1.6~$\\mu$m) linear polarizations for apparent magnitudes from 7th to beyond 14th. This same Galactic region has already been surveyed by $IRAS$ \\citep{NEU84}, NVSS \\citep{NVS98}, $MSX$ \\citep{PRI01}, GLIMPSE \\citep{GLIMP}, the $^{13}$CO Galactic Ring Survey \\citep[GRS;][]{JAC06}, 2MASS \\citep{2MASS}, MIPSGAL \\citep{CAR09}, $WISE$ \\citep{WRI10}, and BGPS \\citep{BCS11}, and so offers excellent data sets for correlative analysis.  The newly available GPIPS background starlight polarimetry, with its roughly one~arcmin stellar sampling, is ideal for exploring the nature of magnetic fields within molecular clouds. A key first step is to measure and characterize the magnetic field within an average molecular cloud that is not engaged in active star formation. With GPIPS polarimetric data, combined with velocity information from GRS, the CF method can be used to create a resolved map of a cloud's embedded plane-of-sky magnetic field strength. Such maps can be analyzed to examine both structural properties and key diagnostic relationships underlying the physical conditions in the cloud.  \\begin{deluxetable*}{cccccc} \\tabletypesize{\\footnotesize} \\centering \\tablewidth{0pt} \\tablecolumns{6} \\tablecaption{Cloud Properties} \\tablehead{\\colhead{Cloud Desig.} & \\colhead{Identifier} & \\colhead{Distance\\,\\tablenotemark{a}} & \\colhead{$V_{LSR}$ Range\\,\\tablenotemark{b}} & \\colhead{Extent} &\\colhead{Mass\\,\\tablenotemark{b}} \\\\ \\colhead{} &\\colhead{} &\\colhead{[kpc]} &\\colhead{[km s$^{-1}$]} &\\colhead{[pc]} &\\colhead{[$M_{\\sun}$]} } \\startdata GRSMC 45.60+0.30 & Cloud 1 & 1.88 & 22.0 - 30.0 & $45 \\times 22$ & $17,000$  \\\\ GRSMC 45.46+0.05 & Cloud 2 & 7.45 & 50.0 - 66.0 & $75 \\times 53$ & $49,000$  \\\\ \\enddata \\label{tab:cloud} \\tablenotetext{a}{\\citet{DUV09}} \\tablenotetext{b}{\\citet{SIM01}} \\end{deluxetable*}  A region near $\\ell=45\\degr.5,b=+0\\degr.2$ was chosen for this quiescent cloud study. In this direction, GRS data reveal the presence of two coincident molecular clouds \\citep{SIM01} separated in velocity, and therefore space. These two clouds are the quiescent cloud GRSMC 45.60+0.3, hereafter `Cloud~1,' and the actively star-forming cloud GRSMC 45.46+0.05, hereafter `Cloud~2' (see Figure \\ref{fig:Cooverlay_paper} top). \\citet{SIM01} and \\citet{DUV09} established general properties for both clouds, as summarized in Table \\ref{tab:cloud}. \\citet{DUV09} used {H\\kern0.1em{\\sc i}} spectral line self-absorption to resolve the near-far kinematic distance ambiguity to place Cloud~1 at a heliocentric distance of 1.88~kpc (near) and Cloud~2 at 7.45~kpc (far).  \\begin{figure} \\plotone{integrated_cropped.eps} \\caption{Map showing extents and locations of Cloud~1 (quiescent) and Cloud~2 (active). Cloud~1 is delineated by a single $^{13}$CO contour representing peak antenna temperature at the 1.0 K level (black contour). Cloud~2 is shown as the filled, gray region, representing $^{13}$CO peak antenna temperature at and above the 6 K level. The bold, dashed curve traces the `spine' of Cloud 1, as discussed in Section 2.3. (\\emph{Bottom}) Gray scale image of Cloud~1 $^{13}$CO integrated line intensity (between $V_{LSR}$ 20 and 32 km s$^{-1}$). Overlaid are the high-quality, 2MASS-matched GPIPS polarization vectors assigned to Cloud~1 (vectors overlapping Cloud~2 have been removed). Vector lengths are proportional to the polarization percentages, with a 2\\% reference vector shown in the lower left corner. Vector orientations show the polarization position angles (directions of the projected maximum of the electric field vector). The GRS data have 45 arcsec resolution, represented by the circle in the upper right.} \\label{fig:Cooverlay_paper} \\end{figure}   This work explores the magnetic field structure toward this region by combining GPIPS starlight polarimetry, GRS $^{13}$CO spectra, and 2MASS photometry. By correlating extinction maps with integrated $^{13}$CO images over each cloud's velocity range, the observed stellar color excess, and therefore background starlight polarimetry, was found to be associated primarily with the quiescent Cloud~1. The CF method was used to estimate the strength of the magnetic field, in the plane of the sky, across the full extent of Cloud~1, with 100~arcsec angular resolution. These results were used to create, for the first time, a resolved map of the magnetic field strength across a molecular cloud and to analyze the contents of this map. In Section~2, the cloud selection and its properties are described. Section~3 discusses the GPIPS observations and Section~4 describes the data analysis. The results are presented in Section~5 and Section~6 contains a discussion of these results. In Section 7, the work is summarized. ",
        "conclusions": "This work presents the first resolved magnetic field map for an entire quiescent molecular cloud, in this case the 40 pc long cloud GRSMC 45.60+0.30. This map has an effective angular resolution of 100 arcsec (0.9 pc), a result made possible by the large number of GPIPS background starlight polarization measurements and coincident $^{13}$CO gas information. Combined with estimates of the cloud thickness and polarization position angle dispersion, the magnetic field was able to be estimated using the CF method for over 900 independent cloud directions. It is with these magnetic field estimates that we next explore relationships between magnetic field strength and gas density as well as between magnetic fields and cloud evolution.  \\subsection{Magnetic Field Strength Dependence on Gas Density}  Within molecular clouds, the magnetic field strength $B$ is correlated with gas density $n_H$. \\citet{TROL8} obtained a $B$-$n_H$ power law slope of 0.4 to 0.6 for their Zeeman work towards dark clouds at moderate gas densities ($> 100$~cm$^{-3}$). \\citet{CRUT9} compiled a list of dark cloud magnetic field measurements and fit a power law index of 0.47$\\pm$0.08 for densities greater than $1000$~cm$^{-3}$. Recently, \\citet{CRU10} completed a comprehensive analysis of Zeeman measurements from the literature, attempting to quantify the actual magnetic field strength probability distribution function (PDF) by comparing the observed line-of-sight magnetic field strength PDF with different models. The {H\\kern0.1em{\\sc i}} Zeeman measurements are particularly relevant because they probe densities similar to the average density found here for Cloud 1. \\citet{CRU10} concluded that the magnetic field is consistent with a two-part model wherein magnetic field strength is independent of n$_{H}$ below 300~cm$^{-3}$ and $log(B) / log(n_H) \\approx 0.65$ above 300~cm$^{-3}$.  The relationship between plane-of-sky magnetic field strength and molecular hydrogen density for Cloud~1 was presented as Fig.~\\ref{fig:bn}. It showed a $B$-$n_{H_2}$ relation power law index of $0.75\\pm0.02$. The Fig.~\\ref{fig:bn} results were combined with the \\citet{CRU10} observational data (gray data and error bars) and their derived model (dashed black line) to become Figure~\\ref{fig:cntb}. The \\citet{CRU10} data are one-dimensional line-of-sight $B_Z$ values. The Cloud~1 data are the two-dimensional plane-of-sky measurements reduced along the y-axis by root two to synthesize one-dimensional $B_X$ (or $B_Y$) values and shifted to positive $n_H$ by 0.3 dex to convert Cloud 1 H$_2$ densities to {H\\kern0.1em{\\sc i}}, as was done by \\citet{CRU10}.  \\begin{figure} \\centering \\includegraphics[angle=90,scale=0.45]{crutcher2010_with_romulan_20120524.eps} \\caption{Log-log plot of one-dimensional magnetic field strength against atomic hydrogen (equivalent) gas volume number density. Line-of-sight $B_Z$ data from \\citet{CRU10} are shown as gray crosses and error bars. The \\citet{CRU10} model is shown as the broken dashed line. Shown as filled diamonds are the plane-of-sky data from Fig.~\\ref{fig:bn}, without error bars, shifted down by 0.15 dex to synthesize one-dimensional $B_X$ values and shifted right by 0.3 dex to count $H_2$ densities twice, as per \\citet{CRU10}. The solid black line is the similarly shifted fit to the plane-of-sky data, with dotted extension beyond the fit region. Note how the new data  depart from the model below $n_H \\sim 300$ cm$^{-3}$ but predict good agreement for higher densities.} \\label{fig:cntb} \\end{figure}  The trend seen for Cloud 1 is inconsistent with the constant magnetic field strength below 300 cm$^{-3}$ in the model of \\citet{CRU10}. Previous work had suggested that below some threshold gas density (between 100 and 300 cm$^{-3}$), the $B$-$n_H$ relationship broke down \\citep{TROL8,CRU10}. The Cloud 1 results presented here show that the $B$-$n_H$ relationship survives to the low densities probed in GRSMC 45.60+0.30, and with an index nearly identical to the one found at higher densities by \\citet{CRU10}.  \\subsection{Cloud and Core Stability and Evolution}  The role of magnetic fields in the support of GRSMC 45.60+0.30 was revealed by the predominance of subcritical values for the mass-to-flux ratios, calculated in Sec. 5.4. In Fig. \\ref{fig:ftm}, many areas of Cloud~1 are magnetically-supported against collapse, including the high-density cores. Importantly, since the mass-to-flux values are upper limits, all subcritical designations are robust.  All of the magnetic cores are subcritical; the magnetic fields likely provide important magnetostatic support. Though the effects of turbulence were not considered, including them would only {\\it reduce} the magnetic support necessary to prevent collapse, making the cores even more stable against gravity.  The magnetic cores are coincident with the high column density regions seen in the $^{13}$CO integrated intensity map (Fig. \\ref{fig:coint}). The arrangement of the magnetic cores along the cloud's spine of more general $^{13}$CO emission and the overall alignment of the spine with the projected magnetic field (Fig. \\ref{fig:paspine}) suggest that the magnetic field may be important in the formation and evolution of both the cloud and the cores. The enhancement of the magnetic field in the high-density regions suggests that the magnetic field has been amplified as the cores formed from more diffuse intercore gas. The average magnetic field strength for the cores is about 1.5 times the mean magnetic field strength for the cloud ($8.3\\pm0.9$ $\\mu$G versus $5.40\\pm0.04$ $\\mu$G). The core mean plane-of-sky magnetic field strength is also higher than the value typically found in the cold HI phase of the ISM \\citep[6~$\\mu$G for the full 3-D field strength;][]{HEIL5}. No correlation was found between gas number density in the cores and their level of mass-to-flux criticality.  These cores would therefore seem to be the best example of magnetic fields, even weak ones, regulating the star formation process. It is noteworthy that 8~$\\mu$G magnetic fields can prevent these $\\sim130$~M$_{\\sun}$ dense cores from collapsing for what must be at least several free-fall times.  \\subsection{Characteristic Core Spacing}  The existence of magnetic cores that are coincident with regions of high $^{13}$CO density and that are relatively uniformly spaced may relate to theories of fragmentation for a self-gravitating fluid cylinder \\citep[``sausage'' instability]{CF53b,N87}. Following the approach of \\citet{JAC10}, Cloud~1 could be considered to be either an incompressible fluid or an infinite isothermal gas cylinder, yielding two distinct characteristic core spacings. For the conditions present in Cloud~1, these spacings are 13 and 17 pc, respectively. They are both larger than the measured mean core spacing of $5.7\\pm0.9$~pc. Because of the low volume densities present in Cloud~1, the gravitationally-driven ``sausage\" instability is unlikely to account for the observed core spacing. The Parker instability \\citep{PAR66} was also considered, but the expected instability scale length is also too large to account for the Cloud's core spacing."
    },
    "1208/1208.3421_arXiv.txt": {
        "abstract": "Interactions between galaxies are very common. There are special kinds of interactions that produce systems called Polar Ring Galaxies (PRGs), composed by a lenticular, elliptical, or spiral host galaxy, surrounded by a ring of stars and gas, orbiting in an approximately polar plane. The present work aims to study \\mbox{AM\\,2020-504}, a PRG with an elliptical host galaxy, and a narrow and well defined ring, probably formed by accretion of material from a donor galaxy, collected by the host galaxy. Our observational study was based on BVRI broad band imagery as well as longslit spectroscopy in the wavelenght range 4100--8600\\AA,  performed at the 1.6\\,m  telescope at the Observat\\'orio do Pico dos Dias (OPD), Brazil. We estimated a redshift of \\textit{z}= 0.01683, corresponding a heliocentric radial velocity of 5045 $\\pm$ 23 km/s. The (B-R) color map shows that the ring is bluer than the host galaxy, indicating that the ring is a younger structure. Standard diagnostic diagrams were used to classify the main ionizing source of selected emission-line regions (nucleus, host galaxy and ring). It turns out that the ring regions are mainly ionized by massive stars while the nucleus presents AGN characteristics. Using two empirical methods,  we found  oxygen abundances  for the H\\,II regions located in the ring in the range 12+log(O/H)=8.3-8.8 dex, the presence of an oxygen gradient across the ring, and that  \\mbox{AM\\,2020-504}  follows the metallicity-luminosity relation of spiral galaxies. These results support the accretion scenario for this object and rules out cold accretion as source for the HI gas in the polar ring. ",
        "introduction": "Galaxies have long been seen as islands distantly scattered and stable in the universe. We now know that galaxies are not randomly distributed in space. They are in groups that are subject to the expansion of the universe and mutual gravitational interaction. The interaction between galaxies has substantially modified the cosmic structures throughout the evolution of the Universe. These events are determined by the attractive character of gravity which in turn induces in larger systems, collisions, tidal forces and dynamical frictions \\citep{1999AJ....117.2695R}. The strong perturbations on the interacting systems are due to the tidal force. This can dismember large quantities of material to form bridges and tails, and thus injecting chemically processed interstellar material into the intergalactic space, contaminating distances up to 10 times larger than the diameter of the iterating galaxies \\citep{1997ASPC..114...71D}.  One of the many types of interactions occur when there is a ring of gas, dust and stars positioned perpendicularly with the galaxy's main plane. These systems are known as polar ring galaxies (PRG), peculiar systems with early-type or elliptical host galaxies.  The term Polar ring galaxies was first introduced by \\cite{1983AJ.....88..909S} and used by Whitmore in the publication \\cite{1987ApJ...314..439W}. \\cite{1990AJ....100.1489W} published his ``Polar Ring Catalog\" (PRC), with a total of 157 objects: 6 kinematically confirmed (rotation, detected in two orthogonal planes), 27 galaxies as ''good candidates\", 73 as ``Possible candidates\", and 51 galaxies as ''related objects\". \\cite{1998ApJ...499..635B} discuss the origin of the fundamental observational properties of polar ring galaxies. \\cite{1998A&AS..129..357F} make a more comprehensive classification of all collisional ring galaxies, which includes the PRGs. Within our neighboring Universe, 20 PRGs have recently been confirmed in the catalog by \\cite{2009Natur.461...43G},  as well as \\cite{2011MNRAS.418..244M} display a new catalogue with candidates to polar-ring galaxies selected from the SDSS.  The \\cite{2003A&A...401..817B} reviewed the two scenario for the formation of PRGs: (1) the fusion that occurs in a frontal collision between two spiral galaxies whose discs are orthogonal, (2) the accretion scenario, in which during the interaction between two galaxies the host collects material from another galaxy to form the ring. Both scenarios require a specific geometric configuration for the formation of a polar ring.  Also, \\cite{2006ApJ...636L..25M} proposed the (3) \\textit{cold accretion scenario} for the formation of isolated PRG's. Based on a large cosmological hydrodynamical simulation, they showed that their formation can occur naturally in a hierarchical universe where most low-mass galaxies are assembled through the accretion of cold gas infalling along megaparsecscale filamentary structures.  Here we report the results of a study of the PRG \\mbox{\\mbox{AM\\,2020-504}}, based on broad band images and long-slit spectroscopy obtained at the Observat\\'orio Pico dos Dias, Brazil. The main goal of this paper is to investigate the scenario of formation  by determinations of the oxygen abundance in the star-forming regions located in the ring and infering  the dust and gas content of the system. This was done through broadband images and spectroscopic data, from which we study the kinematics, surface and aperture photometry. In Section \\ref{sec2} we present a review of \\mbox{\\mbox{AM\\,2020-504}}. Observation and data reductions are presented in Section \\ref{sec3}. The results and discution are presented in Section \\ref{sec4}, while the conclusions are given in Section \\ref{con}. ",
        "conclusions": "\\label{con}  This work presents a study of \\mbox{AM\\,2020-504}, a galaxy with a well defined polar ring surrounding an elliptical host galaxy (\\citealt{1987IAUS..127..413W}, \\citealt{1993A&A...267...21A} and \\citealt{2002A&A...391..103I}). The ring was probably formed by accretion of material from a donnor galaxy during an interaction event. In the field around the galaxy, we did not find any nearby object that might have given material for the formation of the ring, but there is a group of nearby galaxies with similar radial velocities.  We estimated a redshift of \\textit{z}= 0.01683, corresponding to a heliocentric radial velocity of 5045$\\pm$23 km/s, confirming the values found by \\cite{1987IAUS..127..413W} and \\cite{1993A&A...267...21A}. The rotation curve of the ring is symmetrical and well behaved. The last two points each side of the rotation curve suggest that the northern and southern portions of the ring have a difference in rotation velocity of about 60 km/s, but this difference is under the error bars. To a certain degree, asymmetries could be explained if the ring was warped.  We found the (B-R) color index averaged 0.35 and 1.73 for the ring and core of the host galaxy respectively. Thus the ring is bluer than the host galaxy (bulge + nucleus), and that is what we expectif the ring is the result of a recent interaction.  The B-band brightness profile along the minor axis of the galaxy is asymmetric due to the ring. The NW peak is higher and corresponds to the bright spots seen in the images. This morphological feature, as the general S--shaped appearence of the ring are in good agreement with the warped model of the polar ring done by \\cite{1993A&A...267...21A}. The light profile along the host galaxy major axis also looks asymmetric on both sides close to the center. This seems to be due to the presence of dust where the ring passes in front of the galaxy, an indication that the near side of the ring is to the NE of the galaxy.  This system is a harbours an AGN as indicated by some diagnostic diagrams. Using two empirical methods based on the emission-lines easily observable, we found: (i) oxygen abundances  for the H\\,II regions located at the ring  in the range 12+log(O/H)=8.3-8.8 dex with an average value of $8.53\\pm0.11$ dex and (ii) the presence of an oxygen gradient across the ring of about $-0.035$ dex/kpc. We also found that  \\mbox{AM\\,2020-504}  follows the metallicity-luminosity relation of typical spiral galaxies. These results support the accretion scenario for this object and rule out cold accretion."
    },
    "1208/1208.1048_arXiv.txt": {
        "abstract": "A precise and accurate determination of the Hubble constant based on Cepheid variables requires proper characterization of many sources of systematic error. One of these is stellar blending, which biases the measured fluxes of Cepheids and the resulting distance estimates. We study the blending of 149 Cepheid variables in M33 by matching archival {\\it Hubble Space Telescope} data with images obtained at the WIYN 3.5-m telescope, which differ by a factor of 10 in angular resolution.  We find that $55\\pm4$\\% of the Cepheids have no detectable nearby companions that could bias the WIYN $V$-band photometry, while the fraction of Cepheids affected below the 10\\% level is $73\\pm4$\\%. The corresponding values for the $I$ band are $60\\pm4$\\% and $72\\pm4$\\%, respectively. We find no statistically significant difference in blending statistics as a function of period or surface brightness. Additionally, we report all the detected companions within 2$\\arcsec$ of the Cepheids (equivalent to 9 pc at the distance of M33) which may be used to derive empirical blending corrections for Cepheids at larger distances. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.5494_arXiv.txt": {
        "abstract": "We study the convergence properties of our implementation of the {\\it moving punctures} approach at very high resolutions for an equal-mass, nonspinning, black-hole binary. We find convergence of the Hamiltonian constraint on the horizons and the $L_2$ norm of the Hamiltonian constraint in the bulk for sixth- and eighth-order finite difference implementations. The momentum constraint is more sensitive, and its $L_2$ norm shows clear convergence for a system with consistent sixth-order finite differencing, while the momentum and BSSN constraints on the horizons show convergence for both sixth- and eighth-order systems. We analyze the gravitational waveform error from the late inspiral, merger, and ringdown. We find that using several lower-order techniques for increasing the speed of numerical relativity simulations actually lead to apparently nonconvergent errors. Even when using standard high-accuracy techniques, rather than seeing clean convergence, where the waveform phase is a monotonic function of grid resolution, we find that the phase tends to oscillate with resolution, possibly due to stochastic errors induced by grid refinement boundaries. Our results seem to indicate that one can obtain gravitational waveform phases to within $0.05$ rad. (and possibly as small as 0.015 rad.), while the amplitude error can be reduced to $0.1\\%$. We then compare with the waveforms obtained using the CCZ4 formalism. We find that the CCZ4 waveforms have larger truncation errors for a given resolution, but the Richardson extrapolation phase of the CCZ4 and BSSN waveforms agrees to within $0.01$ rad., even during the ringdown. ",
        "introduction": "\\label{sec:Introduction}   Numerical relativity (NR) has progressed rapidly since the breakthroughs of 2005~\\cite{Pretorius:2005gq, Campanelli:2005dd, Baker:2005vv} that  allowed for the long-term evolution of black-hole binaries (BHBs). Among NR's significant achievements are its contributions towards the modeling of astrophysical gravitational wave sources that will be relevant for the first direct detection and parameter estimation by gravitational wave observatories~\\cite{Aylott:2009ya, Aylott:2009tn}.  NR has also made contributions to the modeling of astrophysical sources, notably, the modeling of the recoil kick imparted to the remnant BH from a BHB merger due to unequal masses~\\cite{Herrmann:2006ks, Baker:2006vn, Gonzalez:2006md}, the remarkable discovery of unexpectedly large recoil velocities from the merger of certain spinning BHBs~\\cite{Herrmann:2007ac, Campanelli:2007ew, Campanelli:2007cga, Lousto:2008dn, Pollney:2007ss, Gonzalez:2007hi, Brugmann:2007zj, Choi:2007eu, Baker:2007gi, Schnittman:2007ij, Baker:2008md, Healy:2008js, Herrmann:2007zz, Herrmann:2007ex, Tichy:2007hk, Koppitz:2007ev, Miller:2008en, Lousto:2011kp, Zlochower:2010sn, Lousto:2010xk, Lousto:2011kp, Lousto:2012su}, and the application of the numerical techniques to combined systems of BHs and neutron stars~\\cite{Sekiguchi:2010ja, Etienne:2011ea, Etienne:2008re, Etienne:2007jg, Rezzolla:2011da, Hotokezaka:2011dh, Sekiguchi:2011zd, Foucart:2011mz, Duez:2008rb}. More mathematical aspects of relativity have also recently been investigated, including the evolution of N-black holes~\\cite{Lousto:2007rj, Campanelli:2007ea, Galaviz:2010te}, the exploration of the no-hair theorem~\\cite{Campanelli:2008dv, Owen:2010vw}, and cosmic~\\cite{Campanelli:2006uy} and topological censorship~\\cite{Ponce:2010fq}, as well as BHBs in dimensions higher than four~\\cite{Shibata:2010wz, Zilhao:2010sr, Witek:2010xi}. The current state of the art simulations can simulate  BHBs with mass ratios as small as $q=1/100$~\\cite{Lousto:2010ut, Sperhake:2011ik} and highly spinning BHBs with intrinsic spins $\\alpha=S_H/M_H^2$ up to (at least) $0.97$~\\cite{Lovelace:2010ne, Lovelace:2011nu}.  Currently these runs are very costly and it is hard to foresee the possibility of completely covering the parameter space densely enough for match filtering the data coming from advanced laser interferometric detectors by the time they become operational.  To reduce the computational costs, several low-accuracy approximations are sometimes used. Among them are the techniques introduced in Ref.~\\cite{Brugmann:2008zz} where the number of buffer zones at AMR boundaries is reduced by lowering the order of finite differencing by successive orders near the AMR boundaries, the use of simple interpolations of spectral initial data rather than using the complete spectral expansion~\\cite{Ansorg:2004ds}, and copying the initial data to the two past time levels for use in prolongation at the initial timestep. All of these approximations proved to be useful for numerical simulations, but each one also has the side effect of introducing a (hopefully) small ${\\cal O}(h)$ error.  In this paper, we examine the effects of these approximations by performing high-resolution simulations of equal-mass, nonspinning BHBs, a problem generally considered well under control. We show that a nonconvergent error, that cannot be detected by simple means, is present when these techniques are used together. However, even when low-accuracy approximations are eliminated, an apparently stochastic error in the waveform phase is still present that prevents us from seeing clean convergence of the waveform even at very high resolutions. We estimate this stochastic phase error is controllable  to within NINJA and NRAR accuracy requirements, but does make it very difficult to get an unambiguous measurement of the waveform phase and phase error.  Here we examine in detail the case of a nonspinning equal-mass binary. The idea is that, any issues of accuracy found for these simple systems will only be compounded by the introduction of different mass ratios (which may require more AMR refinement levels and have a lower effective resolution) and spins (which reduces the smoothness of the data and leads to more complicated motion of the binary). ",
        "conclusions": "We analyze the accuracy of a BHB simulation by examining the preservation of the individual horizon masses, convergence of the constraints, and convergence of the magnitude and phase of the $(\\ell=2,m=2)$ mode $\\psi_4$ extracted at $r=50M$. As seen in Fig.~\\ref{fig:hmass}, the fast, but low-accuracy approximations lead to poor conservation of the mass compared to the slower, but more accurate, techniques.  For the high-accuracy techniques, we find convergence of the constraints to the expected order, as seen in Fig.~\\ref{fig:H_pi_nopi}. A residual, possibly stochastic, phase error is seen in the waveform itself. By comparing the waveforms generated using different techniques and different evolution systems, as shown ~in Figs.~\\ref{fig:phase_error}~and~\\ref{fig:phase_v_res_z4_bssn}, we find consistency in the phase at the level of $0.05$ rad. over the entire waveform."
    },
    "1208/1208.0752_arXiv.txt": {
        "abstract": "\\let\\thefootnote\\relax\\footnotetext{michael.koehn@aei.mpg.de,~~~jlehners@aei.mpg.de,~~~ovrut@elcapitan.hep.upenn.edu} It was recently demonstrated that, when coupled to ${\\cal{N}}=1$ supergravity, the Dirac-Born-Infeld (DBI) action constructed from a single chiral superfield has the property that when the higher-derivative terms become important, the potential becomes negative. Thus, DBI inflation cannot occur in its most interesting, relativistic regime. In this paper, it is shown how to overcome this problem by coupling the model to one or more additional chiral supermultiplets. In this way, one obtains effective single real scalar field DBI models with arbitrary positive potentials, as well as multiple real scalar field DBI inflation models with hybrid potentials. \\vspace{.3in} \\noindent ",
        "introduction": "Inflation is a possible solution to the flatness and horizon puzzles of standard big bang cosmology. It was discovered more than 30 years ago \\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi}, and is considered by many as the leading cosmological model of the early universe. This is due in large part to its ability to generate nearly scale-invariant density perturbations at the same time as addressing the above-mentioned puzzles. However, inflation is not unique in this regard. For example, during an ekpyrotic phase \\cite{Khoury:2001wf,Lehners:2007ac,Lehners:2008vx}, where the universe contracts very slowly, the same big bang puzzles can be addressed and nearly scale-invariant density perturbations can be generated\\footnote{In these models, one must also understand the transition from a contracting to an expanding phase--this remains an open issue, but see \\cite{Turok:2004gb,Buchbinder:2007ad,Buchbinder:2007tw,Lehners:2011kr}.}. It follows, therefore, that to understand the actual history of our universe, we must make progress in two directions. On one hand, it is important to work out the detailed predictions of the various models of the early universe--in particular, the different predictions they make regarding the non-Gaussian features in the primordial density perturbations \\cite{Maldacena:2002vr,Chen:2010xka,Buchbinder:2007at,Koyama:2007if,Lehners:2007wc,Lehners:2008my,Lehners:2009ja,Lehners:2009qu}. On the other hand, it is imperative to develop the microphysical structure of the various cosmological models. In this paper, we will be mainly interested in this second aspect--with the focus on inflationary models.  Our aim is to study inflationary theories with higher-derivative kinetic actions in the context of four-dimensional, ${\\cal{N}}=1$ supergravity\\footnote{A study of DBI inflation in global supersymmetry (with an added Einstein-Hilbert term) was performed in \\cite{Sasaki:2012ka}.}. Although our work will be purely within this supergravity context, the motivation stems from string theory. There, the dynamics of D-branes and M5-branes are described by the Dirac-Born-Infeld (DBI) action \\cite{Leigh:1989jq}\\footnote{The effective description in terms of the DBI action is valid at arbitrary velocity, but only as long as the proper acceleration of the branes is small.}. This action is unusual in that it contains higher-derivative terms which are essential to understanding its dynamics\\footnote{Higher-derivative terms involving the extrinsic and intrinsic brane curvatures--such as those discussed in \\cite{Khoury:2012dn,Ovrut:2012wn}--can arise as well. We will not consider these couplings here, but note that they might be significant in certain applications.}. Furthermore, interactions between branes (and anti-branes) can generate an effective potential \\cite{Dvali:1998pa,Lima:2001nh,Buchbinder:2002pr,Kachru:2003sx,Gray:2007zza}. In such a setting, inflationary models based on the DBI action, in which the inflaton field is identified with a position modulus of the brane, have been constructed and shown to lead to interesting observational predictions--such as equilateral non-Gaussianities \\cite{Silverstein:2003hf,Alishahiha:2004eh}. These models have mainly been analyzed in non-supersymmetric effective field theory. However, realistic string compactifications typically preserve minimal supersymmetry in four dimensions--see, for example \\cite{Braun:2005nv,Lukas:1998yy}. It is of interest, therefore, to re-formulate these models within the context of four-dimensional, $\\mac{N}=1$ supergravity.  In a recent paper \\cite{Koehn:2012ar}, we developed a formalism for coupling chiral supermultiplets with higher-derivative kinetic terms to supergravity. Restricting to a single chiral superfield, we constructed a supergravitational generalization of the single real scalar DBI action. This supergravity  theory then contains the DBI action of two real scalar fields--the constituents of the lowest component of the chiral supermultiplet--along with a specific potential energy. In the process, however, we discovered that when the higher-derivative terms become significant, the potential energy necessarily becomes negative--regardless of the form of the superpotential. Thus, with a single chiral supermultiplet, DBI inflation cannot occur! In this paper, we will show how this restriction can be overcome by coupling the supergravity DBI theory to one or more additional chiral superfields--each, however, with canonical two-derivative kinetic terms. Such couplings can lead to positive, inflationary potentials via the elimination of the new auxiliary fields. The required couplings are similar, and in some cases identical, to those previously considered in several two-derivative inflationary models in supergravity \\cite{Kawasaki:2000yn,Kallosh:2010xz,Lazarides:1995vr}. However, in the higher-derivative context, they lead to a number of new features, and to different predictions for cosmological observations.  We have two main results. 1) Within the context of ${\\cal{N}}=1$ supergravity, we provide a method for obtaining DBI inflation for a single real scalar component of a chiral superfield with an arbitrary potential energy. This is accomplished both when the higher-derivative terms are negligible and, more importantly, in the relativistic regime where the higher-derivative terms are dominant. We achieve this by coupling the single chiral superfield DBI theory to one additional chiral supermultiplet--with two-derivative kinetic energy, constrained K\\\"ahler potential and specific holomorphic couplings. 2) We show how one can obtain multi-real-field DBI models with positive potentials. There are two possibilities here. First, within the context of the models just discussed one can allow the scalar superpartner of the inflaton field to fully participate in the dynamics. This is accomplished by easing restrictions on the K\\\"ahler potential. In this case, the potential for the second real scalar field is automatically determined. Second, and more generally, one can couple the supergravity DBI theory to two or more additional chiral supermultiplets--in which case there is more freedom in constructing multi-field potentials. The multi-real-scalar-field models are of clear phenomenological interest, since they can each be compatible with current observational data while making predictions that are testable in the near future \\cite{Langlois:2008qf,Kidani:2012jp}.  The plan of the paper is the following. In Section \\ref{SectionReview}, we review the construction of single chiral superfield DBI actions in ${\\cal{N}}=1$ supergravity. This reveals that, in the relativistic regime, the potential for both real component scalars in the  DBI action is negative and thus prohibits inflation from occurring. In Section \\ref{SectionSingle}, we show how the inclusion of a second chiral supermultiplet modifies this conclusion. In fact, via a judicious choice of both the K\\\"ahler potential and superpotential, this allows arbitrary positive potentials to be constructed for a single real DBI scalar field--while simultaneously fixing the remaining three real scalars. In the beginning of the next section, we briefly discuss how this theory can be modified so that both real component scalars of the DBI superfield become dynamical. In Section \\ref{SectionMulti}, we introduce a third chiral superfield. This allows us to construct a more general class of multi-field models of DBI inflation in supergravity, including, for example, models with inflationary potentials of the hybrid type. We conclude in Section \\ref{SectionConclusions}. ",
        "conclusions": "\\label{SectionConclusions}  One of the most important problems in cosmology is to find a scenario for the early universe that is not only in agreement with observations, but is also rooted in a sensible microphysical theory. Only in this way can cosmology and particle physics be united, and a consistent theory of our universe be obtained. While still far from this goal, we have analyzed a small aspect of the problem in this paper--showing how to construct models of DBI inflation in four-dimensional ${\\cal{N}}=1$ supergravity.  Our recent supergravity analysis of higher-derivative actions showed that if one tries to construct a model of DBI inflation from a  single chiral superfield, it is bound to fail--since the potential becomes negative when the higher-derivative terms become important. In this paper, we circumvented this problem by coupling the theory to one or more additional chiral superfields. In fact, the construction in Section \\ref{SectionMulti} can be generalized to an arbitrary number $N$ of chiral superfields--each with two-derivative kinetic terms and appropriately constrained K\\\"ahler potential--and considering a superpotential of the form \\be W = S w(\\Phi^1,\\Phi^2,\\dots,\\Phi^N) . \\ee Then, not only can the potential energy be positive but one can construct a wide range of potential functions for the original DBI scalar $\\phi = \\sqrt{2} \\Re(\\Phi^1)$ and $N-1$ additional real scalars. The remaining real scalars, that is, the two making up the lowest component of $S$ and one scalar in the lowest component of all the other chiral superfields, can be stabilized with masses above the Hubble scale if the K\\\"{a}hler potential satisfies certain requirements discussed in the text. Our analysis can be viewed as a ``proof-in-principle'' that models of multi-real-scalar-field DBI inflation can be constructed in ${\\cal{N}}=1$ supergravity.  A crucial feature of the analysis of chiral superfields with higher-derivative actions is that, via the elimination of the auxiliary fields, the potential energy  generically depends not only on the superpotential, but on the strength of the higher-derivative terms as well. Thus, in general, the potential changes during the dynamical evolution. In this paper, we have shown that, for the constrained K\\\"ahler potentials and superpotentials above, this turns out not to be the case. The contributions to the potential that depend on the higher-derivative terms vanish in the region of field space of dynamical interest. Thus, the potential remains unchanged as the higher-derivative terms become large or small. This feature considerably simplifies the study of the models considered here, and renders them more accessible for deriving their predictions for cosmological observations. We hope to pursue this topic in the near future.  Our construction illustrates that it is far from straightforward to realize DBI inflation in ${\\cal{N}}=1$ supergravity. We have shown one way in which the desired positive potentials can be obtained from an effective model-building point of view. It is interesting to ask whether there exist other ways of realizing DBI inflation within the context of supergravity. More importantly, however, is the question of whether or not such constructions can be obtained from a full-fledged string compactification, or from some other fundamental theory of particle physics. These are pertinent questions for future research."
    },
    "1208/1208.5800_arXiv.txt": {
        "abstract": "The discovery of large numbers of young low-mass stars and brown dwarfs over the last decade has made it possible to investigate star formation and early evolution in a previously unexplored mass regime. In this review, we begin by describing surveys for low-mass members of nearby associations, open clusters, star-forming regions and the methods used to characterize their stellar properties. We then use observations of these populations to test theories of star formation and evolution at low masses. For comparison to the formation models, we consider the initial mass function, stellar multiplicity, circumstellar disks, protostellar characteristics, and kinematic and spatial distributions at birth for low-mass stars and brown dwarfs. To test the evolutionary models, we focus on measurements of dynamical masses and empirical Hertzsprung-Russell diagrams for young brown dwarfs and planetary companions.  Posted with permission from the Annual Review of Astronomy and Astrophysics, Volume 50 \\copyright 2012 by Annual Reviews, http:// www.annualreviews.org. ",
        "introduction": "Molecular clouds give birth to stars across a wide range of masses. The most massive stars are born on the main sequence while low-mass stars must contract for tens to hundreds of millions of years before becoming hot enough for sustained hydrogen fusion. A great deal of observational and theoretical effort has been invested in understanding star formation and pre-main-sequence evolution and their dependence on stellar mass. Over the last decade, it has become possible to extend these studies to the least massive stars and brown dwarfs as surveys have uncovered them in large numbers and at a variety of ages. In this way, star formation and early evolution can be investigated across more than four orders of magnitude in stellar mass, providing more stringent tests of the theories for these processes.   In this review, we summarize the theoretical and observational work on the formation and early evolution of low-mass stars and brown dwarfs. We define a ``low-mass star\" as having a mass between $\\sim0.2$~$M_\\odot$ and the hydrogen burning mass limit \\citep[$\\sim0.075$~$M_\\odot$,][]{bur97,cha00c}. All free-floating objects (as well as some companions) below this mass range are considered brown dwarfs. Some studies have adopted the deuterium burning limit \\citep[$\\sim0.012$~$M_\\odot$,][]{cha00a,spi11} as the lower limit for the definition of a brown dwarf \\citep{bas00}, but this is not done here since deuterium burning has a negligible impact on stellar structure and evolution \\citep{cha07ppv}. Since we are examining the formation and early evolution of low-mass stars and brown dwarfs, we focus on objects with ages $\\lesssim100$~Myr, which is the time-scale for low-mass stars to approach the main sequence. However, we also consider the properties of older stars and brown dwarfs that can help constrain the formation and evolutionary theories, such as their multiplicity and mass function.   This article complements previous reviews concerning low-mass stars and brown dwarfs. \\citet{bas00} reviewed early discoveries of brown dwarfs and \\citet{kir05} reviewed the spectral classification of (mostly older) L and T dwarfs, whereas we focus on the discovery and characterization of low-mass stars and brown dwarfs at young ages. \\citet{whi07} and \\citet{cha00c} reviewed theories for the formation and evolution of brown dwarfs, respectively, and \\citet{luh07ppv} and \\citet{bur07ppv} compared the predictions of the formation models to observations. We summarize the latest developments in those theories and update the previous reviews of the observational constraints on the formation of brown dwarfs. ",
        "conclusions": "The absence of a significant dependence of the abundance of brown dwarfs on stellar density or the presence of O stars, the ability of brown dwarfs to form in isolation and in wide binaries, and the likely existence of protostellar brown dwarfs together provide compelling evidence that brown dwarfs can form without the involvement of tidal shear in massive cluster-forming cores, dynamical interactions, disk fragmentation around solar-type stars, or photoionizing radiation. Thus, it seems likely that the one remaining formation mechanism from Section~\\ref{sec:theory}, turbulent fragmentation, is responsible for some fraction of brown dwarfs, and perhaps most of the ones in low-density regions like Taurus. It is possible that other proposed mechanisms also produce brown dwarfs, particularly in dense clusters, although there is not any clear observational evidence of this so far. One of the most promising avenues for better understanding the formation of brown dwarfs is continued study of low-mass protostars by identifying them in larger numbers with data from {\\it Spitzer}, WISE, and {\\it Herschel Observatory} and by detailed followup observations with facilities like the Atacama Large Millimeter Array.  Measurements of dynamical masses for a small number of young low-mass stars and brown dwarfs have confirmed the theoretical prediction that the hydrogen burning mass limit occurs near a spectral type of M6 for ages of $\\sim1$--100~Myr.  Based on a wide variety of studies, young objects later than $\\sim$M9, including planetary-mass companions, exhibit unusually red colors and faint absolute magnitudes at near-IR wavelengths relative to older field dwarfs. It appears that this behavior can be explained with model atmospheres that have clouds, low gravities, and non-equilibrium chemistry. Further testing and refinement of the atmospheric and evolutionary models at young ages and low masses will require additional measurements of dynamical masses and larger samples of well-characterized young L and T dwarfs."
    },
    "1208/1208.5031_arXiv.txt": {
        "abstract": "X-ray evidence for ultra-fast outflows (UFOs) has been recently reported in a number of local AGNs through the detection of blue-shifted Fe XXV/XXVI absorption lines. We present the results of a comprehensive spectral analysis of a large sample of 42 local Seyferts and 5 Broad-Line Radio Galaxies (BLRGs) observed with \\emph{XMM-Newton} and \\emph{Suzaku}. We detect UFOs in $\\ga$40\\% of the sources. Their outflow velocities are in the range $\\sim$0.03--0.3c, with a mean value of $\\sim$0.14c. The ionization is high, in the range log$\\xi$$\\sim$3--6~erg~s$^{-1}$~cm, and also the associated column densities are large, in the interval $\\sim$$10^{22}$--$10^{24}$~cm$^{-2}$. Overall, these results point to the presence of highly ionized and massive outflowing material in the innermost regions of AGNs. Their variability and location on sub-pc scales favor a direct association with accretion disk winds/outflows. This also suggests that UFOs may potentially play a significant role in the AGN cosmological feedback besides jets and their study can provide important clues on the connection between accretion disks, winds and jets. ",
        "introduction": "The recent detection of blue-shifted Fe~XXV/XXVI absorption lines in the X-ray spectra of several Seyferts and quasars suggests the presence of highly ionized and mildly-relativistic outflows in the center of these AGNs (e.g., Chartas et al.~2002, 2003; Pounds et al.~2003; Dadina et al.~2005; Markowitz et al.~2006; Braito et al.~2007; Cappi et al.~2009; Reeves et al.~2009; Giustini et al.~2011). They are possibly directly connected with accretion disk winds/outflows and the high associated outflow rate and mechanical power suggest they might be able to provide an important contribution to the expected AGN cosmological feedback (e.g., King 2010; Tombesi et al.~2012). Here we describe the systematic analysis and characterization of these so called Ultra-fast Outflows (UFOs) on a large sample of both radio-quiet and radio-loud AGNs. These results are discussed in detail in Tombesi et al.~(2010a, b; 2011a, b).     \\begin{figure}[!t] \\centering \\includegraphics[width=5.5cm,height=5cm,angle=0]{tombesi_fig1.ps} \\hspace{0.2cm} \\includegraphics[width=5.5cm,height=4cm,angle=0]{tombesi_fig2.ps} \\caption{\\emph{Left panel:} XMM-Newton EPIC-pn spectrum of PG~1211$+$143, best-fit model and background spectrum. \\emph{Right panel:} Example of $\\Delta\\chi^2$ distribution from 1000 Monte Carlo simulations.} \\end{figure} ",
        "conclusions": ""
    },
    "1208/1208.3201_arXiv.txt": {
        "abstract": "Pulsars are rotating neutron stars that are seen to slow down, and the spin-down rate is thought to be due to magnetic dipole radiation\\cite{pacini68,gunnostriker69}. This leads to a prediction for the braking index $n$, which is a combination of the spin period and its first and second time derivatives. However, all observed values\\cite{espinozaetal11} of $n$ are below the predicted value of $3$. Here we provide a simple model that can explain the rotational evolution of young pulsars, including the $n=2.51$ of the 958-year-old pulsar in the Crab nebula\\cite{lyneetal93}. The model is based on a decrease in the effective moment of inertia due to an increase in the fraction of the stellar core that becomes superfluid as the star cools through neutrino emission. The results suggest that future large radio monitoring campaigns of pulsars will yield measurements of the neutron star mass, nuclear equation of state, and superfluid properties. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.3492_arXiv.txt": {
        "abstract": "*{~~~~~ Observations of the spatial distributions of young stars in star-forming regions can be linked to the theory of clustered star formation using spatial statistical methods. The MYStIX project provides rich samples of young stars from the nearest high-mass star-forming regions. Maps of stellar surface density reveal diverse structure and subclustering. Young stellar clusters and subclusters are fit with isothermal spheres and ellipsoids using the Bayesian Information Criterion to estimate the number of subclusters. Clustering is also investigated using Cartwright and Whitworth's $Q$ statistic and the inhomogeneous two-point correlation function. Mass segregation is detected in several cases, in both centrally concentrated and fractally structured star clusters, but a few clusters are not mass segregated.}   \\abstract{~~~~~ Observations of the spatial distributions of young stars in star-forming regions can be linked to the theory of clustered star formation using spatial statistical methods. The MYStIX project provides rich samples of young stars from the nearest high-mass star-forming regions. Maps of stellar surface density reveal diverse structure and subclustering. Young stellar clusters and subclusters are fit with isothermal spheres and ellipsoids using the Bayesian Information Criterion to estimate the number of subclusters. Clustering is also investigated using Cartwright and Whitworth's $Q$ statistic and the inhomogeneous two-point correlation function. Mass segregation is detected in several cases, in both centrally concentrated and fractally structured star clusters, but a few clusters are not mass segregated.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.2341_arXiv.txt": {
        "abstract": "We explore the effects of nuclear masses on the temperature and neutron density conditions required for $r$-process nucleosynthesis using four nuclear mass models augmented by the latest atomic mass evaluation. For each model we derive the conditions for producing the observed abundance peaks at mass numbers $A\\sim 80$, 130, and 195 under the waiting-point approximation and further determine the sets of conditions that can best reproduce the $r$-process abundance patterns ($r$-patterns) inferred for the solar system and observed in metal-poor stars of the Milky Way halo. In broad agreement with previous studies, we find that (1) the conditions for producing abundance peaks at $A\\sim 80$ and 195 tend to be very different, which suggests that, at least for some nuclear mass models, these two peaks are not produced simultaneously; (2) the typical conditions required by the critical waiting-point (CWP) nuclei with the $N = 126$ closed neutron shell overlap significantly with those required by the $N=82$ CWP nuclei, which enables coproduction of abundance peaks at $A\\sim 130$ and 195 in accordance with observations of many metal-poor stars; and (3) the typical conditions required by the $N = 82$ CWP nuclei can reproduce the $r$-pattern observed in the metal-poor star HD~122563, which differs greatly from the solar $r$-pattern. We also examine how nuclear mass uncertainties affect the conditions required for the $r$-process and identify some key nuclei including $^{76}$Ni to $^{78}$Ni, $^{82}$Zn, $^{131}$Cd, and $^{132}$Cd for precise mass measurements at rare-isotope beam facilities. ",
        "introduction": "\\label{sec:introduction} Nucleosynthesis via rapid neutron capture, the $r$-process, is a major mechanism for producing the elements heavier than Fe~\\cite{Burbidge1957,Cameron1957}. Understanding this process requires knowledge of properties such as masses, $\\beta$-decay lifetimes, and neutron-capture cross sections for a large number of extremely neutron-rich nuclei far from stability (e.g., ~\\cite{Cowan1991,Qian2003,Arnould2007}). Most of this nuclear input is beyond the reach of experiments in the foreseeable future and, therefore, must be calculated with guidance from existing data and from measurements to be made at rare-isotope beam facilities such as the GSI Facility for Antiproton and Ion Research (FAIR), the Facility for Rare Isotope Beams (FRIB), the Heavy Ion Research Facility in Lanzhou Cooling Storage Ring (HIRFL-CSR), and Rikagaku Kenkyusho (RIKEN, Institute of Physical and Chemical Research, Japan). In this paper we explore the importance of nuclear masses in estimating the temperature and neutron density conditions required for a specific $r$-process scenario, where neutron-capture reactions are in equilibrium with the reverse photodisintegration reactions; i.e., there is $(n,\\gamma)\\rightleftharpoons(\\gamma,n)$ equilibrium (e.g.,~\\cite{Cowan1991,Qian2003,Arnould2007}). Using four nuclear mass models, we show that the required conditions can be determined mostly from the neutron-separation energies for a small number of critical nuclei with $N=50$, 82, and 126 closed neutron shells. For each model, we further determine the best-fit sets of conditions to reproduce the $r$-process abundance pattern ($r$-pattern) inferred for the solar system and observed in metal-poor stars of the Milky Way halo. This allows us to draw several interesting conclusions regarding the production of different parts of the overall $r$-pattern in an $(n,\\gamma)\\rightleftharpoons(\\gamma,n)$ equilibrium scenario. We also illustrate the effects of nuclear mass uncertainties on the required $r$-process conditions and identify the key nuclei that have the largest impact and, therefore, are important candidates for precise mass measurements at rare-isotope beam facilities.  We first give a brief overview of the $r$-process. Detailed reviews can be found in Ref.~\\cite{Cowan1991,Qian2003,Arnould2007}. Historically, the abundance distribution of nuclei in the solar system played a crucial role in studies on the origin of the elements~\\cite{Burbidge1957,Cameron1957}. One of the prominent features of this distribution is the presence of three sets of double peaks in the region beyond the Fe group nuclei. This was recognized as signatures of two distinct processes of neutron capture: a slow ($s$) one encountering the $N=50$, 82, and 126 closed neutron shells in the stable region and a rapid ($r$) one encountering the same in the extremely neutron-rich region of the nuclear chart~\\cite{Burbidge1957,Cameron1957}. Specifically, the peaks at mass numbers $A\\sim 80$, 130, and 195 were produced by the $r$-process and represent the crucial features of the solar $r$-pattern, which is derived by subtracting the $s$-process contributions from the net solar abundances (e.g.,~\\cite{Kappeler2011}).  In order to fully understand the $r$-process, we need conditions such as temperature and neutron density in the associated astrophysical environments in addition to the properties of a large number of extremely neutron-rich nuclei. Neither the astrophysical nor the nuclear input is firmly established, although much progress has been made over the past two decades~\\cite{Cowan1991,Qian2003,Arnould2007}. Proposed astrophysical sites for the $r$-process include neutrino-driven winds from proto-neutron stars formed in core-collapse supernovae (CCSNe)~\\cite{Woosley1992,Meyer1992,Woosley1994}, shocked surface layers of O-Ne-Mg cores associated with low-mass CCSNe~\\cite{Ning2007}, winds from accretion disks of black holes formed in high-mass CCSNe~\\cite{Pruet2003,Surman2006,Surman2008,Wanajo2012}, He shells of metal-poor CCSNe~\\cite{Epstein1988,Banerjee2011}, and ejecta from neutron star mergers~\\cite{Lattimer1977,Freiburghaus1999a,Goriely2011,Korobkin2012}. There are large uncertainties in the conditions associated with all CCSNe environments due to the substantial uncertainties in modeling such environments (e.g.,~\\cite{Janka2012}), especially when neutrino transport in hot and dense nuclear matter is considered (e.g.,~\\cite{Martinez2012,Roberts2012}). While recent studies lend much support to neutron star mergers being an $r$-process site~\\cite{Goriely2011,Korobkin2012}, it remains to be seen whether such models are consistent with the history of $r$-process enrichments in the Milky Way and in its satellite dwarf galaxies (e.g.,~\\cite{Qian2000,Argast2004,Donder2004}). Further, how sensitive these models are to the uncertainties in the current understanding of the nuclear equation of state (e.g.,~\\cite{Page2006}) and to the numerical treatment of the merger dynamics remains to be studied in detail.  The conditions in the astrophysical environments relevant for the $r$-process ultimately boil down to the seed nuclei for neutron capture at the beginning of the process and the temperature $T(t)$ and the neutron (number) density $n_n(t)$ as functions of time $t$ during the process (e.g.,~\\cite{Freiburghaus1999b}). In cases where neutrino interactions are important, the time evolution of neutrino fluxes and energy spectra is also required (e.g.,~\\cite{Epstein1988,Banerjee2011}). In the rest of the paper we will ignore neutrinos and focus on a broad class of astrophysical environments where matter undergoes $r$-processing at $T\\gtrsim 10^9$~K and $n_n\\gtrsim 10^{20}$~cm$^{-3}$. For such high temperatures and neutron densities, previous studies have shown that $(n,\\gamma)\\rightleftharpoons(\\gamma,n)$ equilibrium is achieved (e.g.,~\\cite{Goriely1996}). In this equilibrium, the abundance distribution in each isotopic chain at a specific proton number $Z$ is almost always strongly peaked at one nucleus. This is referred to as a waiting-point (WP) nucleus because, upon reaching it, the $r$-process must wait for it to $\\beta$-decay before producing heavier nuclei. Under this so-called WP approximation, the $r$-process path is defined by all the WP nuclei heavier than the seed nuclei, and the progress along this path is regulated by the $\\beta$-decay of these WP nuclei. So long as this approximation is valid, there is no need to follow neutron capture and photodisintegration reactions, which greatly simplifies the $r$-process calculation.  Due to the equilibrium between neutron capture and photodisintegration reactions, the abundance ratio between two neighboring isotopes is given by the Saha equation (e.g.,~\\cite{Cowan1991,Qian2003,Arnould2007}): \\begin{equation}\\label{equilibrium} \\frac{Y(Z,A+1)}{Y(Z,A)}=n_n\\left(\\frac{2\\pi\\hbar^2}{m_u k T}\\right)^{\\frac{3}{2}} \\frac{G(Z,A+1)}{2G(Z,A)}\\left(\\frac{A+1}{A}\\right)^{\\frac{3}{2}} \\exp\\left[\\frac{S_n(Z, A+1)}{ k T}\\right], \\end{equation} where $\\hbar$ is the Planck constant, $m_u$ is the atomic mass unit, $k$ is the Boltzmann constant, $(Z,A)$ indicates a nucleus with proton number $Z$ and mass number $A$, and $Y$, $G$, and $S_n$ denote the number abundance, partition function, and neutron separation energy of the appropriate nucleus, respectively. For a specific isotopic chain, the corresponding WP nucleus has the largest abundance and is determined by the partition functions and neutron separation energies of the relevant nuclei for fixed $T$ and $n_n$. As can be seen from the exponential dependence on the neutron separation energy in Eq.~(\\ref{equilibrium}), nuclear masses are among the most important input for modeling the $r$-process. The other crucial input is $\\beta$-decay lifetimes of the relevant nuclei.  Over the past two decades, tremendous progress has been made in measuring nuclear properties relevant for the $r$-process. For example, the $\\beta$-decay half-lives of 38 very neutron-rich isotopes bordering the $r$-process path have been measured recently~\\cite{Nishimura2011}. In addition, the masses of a group of nuclei including $^{80}$Zn~\\cite{Baruah2008, Sun2008b} and $^{130}$Cd~\\cite{Dillmann2003} have been measured with a very high accuracy~\\cite{Audi2011}. Meanwhile, considerable advance has been made in the theoretical investigation of nuclear masses. The four nuclear mass models used in this paper span from the macroscopic-microscopic kind, represented by the finite-range droplet model (FRDM)~\\cite{Moller1995} and a more recent Weizs\\\"acker-Skyrme (WS*) model~\\cite{Wang2010}, to the microscopic kind, represented by the Skyrme-Hartree-Fock-Bogolyubov mean-field (HFB-17) model~\\cite{Goriely2009} and the relativistic mean-field (RMF) model~\\cite{Geng2005}. These models can reproduce the experimentally known neutron separation energies with a root-mean-square (rms) deviation of 0.399 (FRDM), 0.332 (WS*), 0.506 (HFB-17), and 0.653 (RMF) MeV, respectively.  Based on the above overview, there are two frontiers of $r$-process research: one focusing on the search for the astrophysical sites and quantification of the conditions therein and the other on acquiring a reliable database for the relevant nuclear input. Observations of elemental abundances in metal-poor stars of the Milky Way halo (see~\\cite{Sneden2008} for a review) have shed important light on the $r$-process sites (e.g.,~\\cite{Qian2007,Farouqi2009}). The observed $r$-patterns also provide an important test of the basic soundness of the nuclear input (e.g.,~\\cite{Farouqi2010}). Of course, the astrophysical and the nuclear input must be coupled together in order to produce an $r$-pattern for comparison with observations. With substantial uncertainties in the current understanding of both the $r$-process sites and the nuclear input, parametrization of the astrophysical conditions is often used in exploring the effects of nuclear input on $r$-process production (e.g.,~\\cite{Freiburghaus1999b,Farouqi2010}). As a practical matter, we adopt the classical approach of using $T$, $n_n$, and the corresponding neutron irradiation time $\\tau$ along with the WP approximation (e.g.,~\\cite{Kratz1993}) to carry out our $r$-process calculations below. Our main purpose is to explore the effects of the four nuclear mass models mentioned above on the $T$ and $n_n$ conditions required for $r$-process nucleosynthesis.  A number of other studies on how the nuclear input impacts $r$-process nucleosynthesis have been carried out recently. The influence of nuclear properties from different mass models on the final $r$-pattern was analyzed in Ref.~\\cite{Farouqi2010}. The effect of neutron capture rates for nuclei near the $A\\sim 130$ peak on the overall $r$-pattern was investigated in Ref.~\\cite{Surman2009}. The sensitivity of the calculated $r$-pattern to the combined effects of the long-term dynamic evolution of the astrophysical environment and the nuclear input was explored in Ref.~\\cite{Arcones2011a}. The effect of long-range correlations for nuclear masses on the production of nuclei immediately below and in the $A\\sim 195$ peak was studied in Ref.~\\cite{Arcones2012}. In all of the above studies parametric astrophysical models that allow $T$ and $n_n$ to evolve with time were adopted. These models are more realistic but the choice of parameters is not so straightforward. In the future, we plan to use similar models to study the interplay between the astrophysical and the nuclear input during the $r$-process. Our goal here is to explore the $T$ and $n_n$ conditions required for the $r$-process in an $(n,\\gamma)\\rightleftharpoons(\\gamma,n)$ equilibrium scenario and the effects of nuclear mass models on these conditions. For this limited goal, we feel that the classical approach based on the WP approximation is adequate.  We give a detailed discussion of the WP approximation in Sec.~\\ref{sec:model}. Using this approximation along with four nuclear mass models, we derive for each model the $T$ and $n_n$ conditions that are required for producing the abundance peaks at $A\\sim 80$, 130, and 195 as observed in the solar system. In Sec.~\\ref{sec:patterns} we describe the classical approach to simulate the $r$-process and use this approach to determine the sets of conditions that can best reproduce the solar $r$-pattern and the $r$-patterns observed in metal-poor stars for each of the adopted nuclear mass models. We discuss our results and give conclusions in Sec.~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss} We have explored the effects of four nuclear mass models (FRDM, WS*, HBF-17, and RMF) on the conditions required by $r$-process nucleosynthesis under the WP approximation. As discussed in Sec.~\\ref{sec:model}, the required $T_9$-$n_n$ conditions are mostly determined by the two-neutron separation energies of the CWP nuclei with $N=50$, 82, and 126 and of those nuclei around them. Figure~\\ref{figModel50} shows the dramatic effect of using the tabulated values in the latest atomic mass evaluation AME2011-preview when they are available to replace the masses predicted by models. As noted in Sec.~\\ref{sec:cond}, the tabulated masses of $^{76}$Ni to $^{78}$Ni and $^{78}$Zn to $^{82}$Zn play crucial roles in determining the conditions required by the $N=50$ CWP nuclei. However, the tabulated masses of $^{76}$Ni to $^{78}$Ni and $^{82}$Zn are extrapolated rather than measured. To emphasize the effects of these masses on the conditions required by the $N=50$ CWP nuclei, we first repeat the calculations of Sec.~\\ref{sec:model} by varying the neutron separation energy of $^{78}$Ni within the estimated uncertainty of 0.946~MeV \\cite{Audi2011} while keeping the other input the same as for Fig.~\\ref{figModel50}(b). The results are shown in Fig.~\\ref{figErr}(a). In comparison with Fig.~\\ref{figModel50}(b), the lower bound on the region of the required $T_9$-$n_n$ conditions stays the same because this is determined by the two-neutron separation energy of $^{82}$Zn (see Sec.~\\ref{sec:cond}), which is not changed. Increasing the neutron separation energy of $^{78}$Ni by 0.946~MeV raises the upper bound from the solid curve [upper bound in Fig.~\\ref{figModel50}(b)] to the dashed curve and decreasing this quantity by the same amount lowers it to the dotted curve. We then repeat the same calculations but vary the neutron separation energy of $^{82}$Zn within the estimated uncertainty of 0.401~MeV~\\cite{Audi2011}. The effects on the lower bound on the region of the required $T_9$-$n_n$ conditions are shown in Fig.~\\ref{figErr}(b). Note that if the neutron separation energy of $^{78}$Ni were lower than its tabulated value by 0.946~MeV while that of $^{82}$Zn were higher by 0.401~MeV, then it would be almost impossible to find any $T_9$-$n_n$ conditions to accommodate all the $N=50$ CWP nuclei. In any case, the significant effects of uncertainties in neutron separation energies on the required $T_9$-$n_n$ conditions shown in panels (a) and (b) of Fig.~\\ref{figErr} clearly demonstrate the importance of precise mass measurements for $^{76}$Ni to $^{78}$Ni and $^{82}$Zn.  \\begin{figure}[h] \\centerline{ \\includegraphics[scale=0.48,angle=0]{fig9.eps} } \\caption{(Color online) Effects of the uncertainty in the neutron separation energy $S_n$ for (a) $^{78}$Ni, (b) $^{82}$Zn, (c) $^{191}$Tb, and (d) $^{197}$Tm on the required $T_9$-$n_n$ conditions. The solid curves in panels (a) and (b) are the same as those in Fig.~\\ref{figModel50}(b). The solid curves in panels (c) and (d) are the same as the dashed curves (for the WS* model) in Fig.~\\ref{figModelall}(d). The shaded regions in each panel show the effects on the required $T_9$-$n_n$ conditions when the corresponding $S_n$ values are varied within the estimated uncertainties. See text for details.} \\label{figErr} \\end{figure}  As in the case of $N=50$ CWP nuclei, we have also made a careful scan of the nuclear chart around the $N = 82$ CWP nuclei and explored the effects of those nuclei with experimentally unknown or poorly measured masses on the $T_9$-$n_n$ conditions required by the $N = 82$ CWP nuclei. When nuclear masses are not known experimentally, we have used the extrapolated masses and uncertainties as these have been proven to have a better predictive power than all available models~\\cite{Fu2011,Sun2011}. We have identified $^{131}$Cd and $^{132}$Cd as additional key nuclei for precise mass measurements.  As can be seen from Fig.~\\ref{figModelall}, the $T_9$-$n_n$ conditions required by the $N=126$ CWP nuclei depend strongly on the nuclear mass model. To assess the impact of uncertainties in nuclear mass models, we varied the neutron separation energies $S_n$ for the relevant nuclei within the known errors or the rms deviations of model predictions for the known masses. Using the WS* model as an example, we show the effects of uncertainties in $S_n$ for $^{191}$Tb and $^{197}$Tm on the required $T_9$-$n_n$ conditions in panels (c) and (d) of Fig.~\\ref{figErr}, respectively. It can be seen that the upper bound on these conditions changes very little when $S_n(^{191}{\\rm Tb})$ is varied within 0.332 MeV (1 rms deviation for the WS* model) but the lower bound is more sensitive to the same variation of $S_n(^{197}{\\rm Tm})$. However, the conditions required by the $N=50$ [Figs.~\\ref{figErr}(a) and 9(b)] and 126 [Figs.~\\ref{figErr}(c) and 9(d)] CWP nuclei do not appear to overlap for the WS* model even when uncertainties in nuclear masses are considered. This suggests that, at least for this model, the $A\\sim 80$ and 195 peaks in the $r$-pattern are most likely produced under very different conditions. Precise mass measurements and better calibrated mass models are needed to make this result more robust.  In conclusion, we have estimated the temperature and neutron density conditions required for $r$-process nucleosynthesis under the WP approximation using four nuclear mass models augmented by the latest atomic mass evaluation AME2011-preview. We have shown that these conditions are mostly determined by the two-neutron separation energies of the $N=50$, 82, and 126 CWP nuclei and those around them. We have also identified some key nuclei including $^{76}$Ni to $^{78}$Ni, $^{82}$Zn, $^{131}$Cd, and $^{132}$Cd for precise mass measurements at rare-isotope beam facilities.  Based on the typical conditions required by the $N=50$, 82, and 126 CWP nuclei shown in Fig.~\\ref{figModelall}, we have performed $r$-process calculations in the classical approach to reproduce the $r$-pattern inferred for the solar system and those observed in metal-poor stars of the Milky Way halo. We have found that (1) at least for the nuclear mass models considered here, the conditions required to produce the peak at $A\\sim 80$ differ greatly from those required to produce the solar $r$-pattern from the peak at $A\\sim 130$ to that at $A\\sim 195$, which reflects that the $T_9$-$n_n$ conditions required by the $N=50$ and 126 CWP nuclei are very different (especially for the WS* model); (2) the solar $r$-pattern from the peak at $A\\sim 130$ to that at $A\\sim 195$, which also closely describes the $r$-patterns in many metal-poor stars, can be reproduced under the conditions required by the $N=126$ CWP nuclei, which has significant overlap with those required by the $N=82$ CWP nuclei, thereby enabling coproduction of the peaks at $A\\sim 130$ and 195; (3) it is plausible to explain the overall $r$-patterns in metal-poor $r$-II stars with a superposition of two sets of $r$-process conditions required by the $N = 50$ and 126 CWP nuclei, respectively; and (4) the non-solar-like $r$-pattern observed in metal-poor stars like HD~122563 can be accounted for by the $r$-process conditions required by the $N = 82$ CWP nuclei. We note that similar results were also obtained by other earlier studies (e.g.,~\\cite{Kratz2007,Schatz2002}).  We recognize that the classical $r$-process approach leaves out many important details, such as the time evolution of temperature and neutron density, the finite duration of the freeze-out, and the breakdown of $(n,\\gamma)\\rightleftharpoons(\\gamma,n)$ equilibrium during the freeze-out. We note that the impact of the details of the freeze-out on the final $r$-pattern~\\cite{Surman2009,Arcones2011a}, especially the formation of the rare-earth peak~\\cite{Mumpower2012a,Mumpower2012b}, has been investigated extensively in other recent studies. However, so long as $(n,\\gamma)\\rightleftharpoons(\\gamma,n)$ equilibrium can be achieved in an $r$-process environment, the conditions immediately before the freeze-out in that environment should be close to those derived here. We intend to carry out parametric studies of the $r$-process based on more detailed and more realistic astrophysical models in the future, and will explore the effects of various nuclear input on such models."
    },
    "1208/1208.5928_arXiv.txt": {
        "abstract": "{} {We present a detailed study of the supernova remnant (SNR) \\snr~in the 0.2-12 keV X-ray band.} {Using data from XMM-Newton we performed a spectro-imaging analysis of \\snr~in order to deduce the basic parameters of the remnant and to search for evidence of a young neutron star associated with it.} {In X-rays the remnant is characterized by a bright arc located in the south-west direction. Its X-ray spectrum can be best described by an absorbed non-equilibrium collisional plasma model with a column density of $N_\\mathrm{H}=1.24_{-0.05}^{+0.07} \\times 10^{22}$ cm$^{-2}$ and a plasma temperature of $6.2^{+0.9}_{-0.8}$ million Kelvin. The analysis indicates a remnant age of 5800 to 7600 years and a distance of $9.8_{-0.7}^{+1.1}$ kpc. The latter suggests a spatial connection with a close-by HII region. We did not find evidence for the existence of a young neutron star associated with the remnant.} {} ",
        "introduction": "Every year, several hundred supernova (SN) events are observed, representing the end state of stellar evolution in which a core-collapse of a massive star or the thermonuclear disruption of a white dwarf takes place.  SNe are the most energetic events which can be observed in the universe. They are often as bright as a whole galaxy. Their extreme brightness allows to see them up to distances of Gpc \\citep{2011AAS...21821901R}. At this scale, though, the only information obtainable from them is the characteristic rising and fading of their light as a function of time, i.e.~their photometric light-curve, and their spectral evolution, both of which are important for their classification \\citep{2008AJ....135..348S}. Most of these events are discovered in the optical band by comparing observations which were taken at different epochs.  In our own Galaxy the rate of observed SNe is small. For core-collapse supernovae it averages to only about two per century. This estimate is suggested by the $\\gamma$-ray radiation from ${}^{26}$Al in the Galaxy \\citep{2006Natur.439...45D} in which a certain yield is expected to be formed in each core-collapse SN. Many SNe, though, remain unobserved due to optical light extinction. The last recorded SN in our Galaxy, the Kepler SN, was observed in AD 1604 and not more than six other SNe have been detected in the Galaxy in the past two thousand years \\citep{2003LNP...598....7G}. The optical light from the two youngest Galactic SNe, G1.9+0.3 and Cas A, were not observed about 100 and 350 years ago \\citep{2008ApJ...680L..41R, 2003LNP...598....7G}.  If the direct light from a SN was missed, for several nearby ones we have the chance to study at least the light of their remnant, which remains visible in various spectral bands for up to $10^5$ years. Although the light echo from the SN Cas A was found recently \\citep{2008ApJ...681L..81R}, it is one of the few cases so far where the SN light could be studied a few hundred years after the SN. The most promising way to learn about the evolution of a SN shock front and the feedback on the evolution of their host galaxy is therefore to study the diffuse emission of SNRs.  In the last years several new SNRs have been detected thanks to the increasing sensitivity of modern X-ray observatories. One of those remnants is \\snr, which was first detected in X-rays by \\citet{2002ASPC..271..391S} and later identified to be a SNR by \\citet{2003PhDTSchaudel} and \\citet{2012MNRAS.419.2623R}. It is a shell-like SNR from which X-ray and radio emission has been detected. In addition, filaments in the infrared and H$_\\alpha$ band were detected in the near proximity of the source \\citep{2012MNRAS.419.2623R}. Whether they are associated with the remnant still remain to be shown.  A first detailed study of the SNR \\snr~was presented by \\citet{2012MNRAS.419.2623R} using ROSAT PSPC data. As this detector had roughly five independent energy channels in the range 0.1-2 keV, the spectral results deduced in their analysis were very limited. They suggested that the X-rays were emitted from a thermal plasma. \\citet{2012MNRAS.419.2623R} therefore were not able to put any constraints with high confidence on the derived spectral parameters. The ROSAT data did not allow them to deduce parameters like the age or the expansion velocity of the remnant.  In this work we report on an XMM-Newton observation which was targeted on SNR \\snr. The results of the spatial and spectral analysis of this data is presented in Section \\ref{sec:data}. A discussion is given in Section \\ref{sec:discussion} in which we use the inferred spectral parameters of \\snr~to derive an estimate for its age, its radius, its expansion velocity and its distance. Section \\ref{sec:summary} provides the concluding summary. ",
        "conclusions": "\\label{sec:summary}  The remnant is characterized by a bright arc in the south-west direction and by diffuse emission with low surface brightness in its western part. We showed that the X-ray emission of \\snr~is in agreement with coming from a collisionally heated plasma which has not reached equilibrium yet. The Sedov analysis leads to the conclusion that the SNR is about 6600 years old and is expanding with a velocity on the order of $\\approx 720$ km/s.  The close-by HII region G296.593--0.975 has a velocity of $+25\\pm1 \\text{ km s}^{-1}$ based on measured hydrogen recombination lines \\citep{1987A&A...171..261C}. With the standard IAU parameter for the distance to the center of our Galaxy $R_0=8.5$ kpc and the solar orbit velocity of $V_0=220$ km/s derived by \\citet{1986MNRAS.221.1023K} and the Galactic rotation model of \\citet{1989ApJ...342..272F} we deduce a distance to the HII region of $9.3\\pm0.6$ kpc. For the error estimate we assumed an uncertainty in the velocity-to-distance conversion of 7 km/s \\citep[e.g.,][]{1988ApJ...333..332C}.  The deduced distance $d_\\text{Sedov}$ is in good agreement with the distance of the close-by HII region G296.593--0.975. This is a strong indicator for a spatial connection between the SNR and the HII region as already suggested by \\citet{2012MNRAS.419.2623R}.  The observation used in this analysis was strongly affected by particle background radiation, which led to a net observation time that was shorter by a factor of three than the approved exposure time. Therefore, only limited statements can be made about the existence of a compact source located near the center of the supernova remnant. Deeper observations might help to clarify this question and the one of the type of the SN in more detail."
    },
    "1208/1208.2025_arXiv.txt": {
        "abstract": "Young radio galaxies (YRGs) provide an ideal laboratory to explore the connection between accretion disk and radio jet thanks to their recent jet formation. We investigate the relationship between the emission-line properties, the black hole accretion rate, and the radio properties using a sample of 34 low-redshift ($z < 0.4$) YRGs. We classify YRGs as high-excitation galaxies (HEGs) and low-excitation galaxies (LEGs) based on the flux ratio of high-ionization to low-ionization emission lines. Using the \\Ha\\ luminosities as a proxy of accretion rate, we find that HEGs in YRGs have $\\sim1$ dex higher Eddington ratios than LEGs in YRGs, suggesting that HEGs have higher mass accretion rate or higher radiative efficiency than LEGs. In agreement with previous studies, we find that the luminosities of emission lines, in particular \\Ha, are correlated with radio core luminosity, suggesting that accretion and young radio activities are fundamentally connected. ",
        "introduction": "Active galactic nuclei (AGN)  play an important role in galaxy evolution by feeding their host galaxies with radiative and/or kinetic energy, leading to the observed scaling relations between black hole (BH) mass and galaxy properties (e.g., Ferrarese \\& Merritt 2000; G{\\\"u}ltekin \\etal~2009; Woo \\etal~2010). Thus, investigating the radiative and kinetic processes of AGN is of fundamental importance for understanding AGN physics as well as feedback mechanisms.  The formation of relativistic jets and its connection to accretion disk remain as an open issue in AGN physics (e.g., Rees 1984; Meier 2003; McKinney 2006; Komissarov \\etal~2007; McKinney \\etal~2012). Thanks to their short dynamical timescale, X-ray binaries in various states have been observed in detail, revealing that individual sources occupy a particular accretion state with various X-ray and radio luminosities (e.g., Fender \\etal~2004; Remillard \\& McClintock 2006). By analogy to X-ray binaries, the power of AGN jets may also depend on the physical states of accretion disk.  However, the disk-jet connection is more complicated in AGNs. It has been known that the radio power is correlated with narrow emission line luminosities, i.e., \\oiii, which is a proxy for the accretion power (e.g., Baum \\& Heckman 1989; Rawlings \\& Saunders 1991), indicating that the jet launching mechanism is connected with accretion disk. However, the correlation shows a large scatter (e.g., Morganti \\etal~1992; Labiano 2008), implying that the nature of the disk-jet connection is complicated and that other physical parameters are necessary to constrain.  The properties of the accretion disk seem to play an additional role in the disk-jet connection. At a given radio luminosity, high-excitation galaxies (HEG) classified with high \\oiii/\\Ha\\ ratio (e.g., Laing \\etal~1994) have an order of magnitude higher \\oiii\\ luminosity than low-excitation galaxies (LEG). Using a sample of 3C radio galaxies, Buttiglione \\etal~(2010) showed that there are two sequences in the radio--emission line luminosity plane, suggesting that systematically different accretion rates or accretion modes are responsible for the separation between HEGs and LEGs.  One of the limiting factors in interpreting the disk-jet connection in powerful radio galaxies is that the lifetime of large radio sources is much longer than the transition timescale of the physical states in the accretion disk (O'Dea \\etal~2009). Thus, comparing radio and accretion powers in radio galaxies with large-scale jets may suffer systematic uncertainties, leading to large scatters in the correlation between optical and radio properties (see Punsly \\& Zhang 2011).  In contrast, compact radio galaxies with small jets triggered by recent activity are very useful to investigate the disk-jet connection, since radio and disk activities are contemporaneous. The ages of compact radio galaxies are typically estimated as $t_{\\rm age} =10^{2}-10^{3}$ years  (e.g., Orienti \\etal~2007; Fanti 2009; Giroletti \\& Polatidis 2009), while the lifetime of extended (up to a few hundred kpc) radio sources are about $10^{6}-10^{7}$ years (e.g., Alexander \\& Leahy 1987; Carilli \\etal~1991; Fanti \\etal~1995; O'Dea \\etal~2009).  Recently, a substantial amount of such compact radio galaxies have been detected and classified with various characteristics: compact symmetric objects (CSO) with a linear scale $\\lesssim 1$ kpc, gigahertz-peaked spectrum (GPS; $\\lesssim 1$ kpc) sources, medium-size symmetric objects ($1-10$ kpc), and compact steep-spectrum (CSS;  $\\lesssim 20$ kpc) sources (See O'Dea \\etal~2009). These young radio galaxies (YRGs) are a good laboratory to investigate the physical link between AGN jets and accretion disks.  YRGs show a close connection between emission-line gas and radio properties. For example, it has been reported that emission-line gas (e.g.,\\oii, \\oiii) is well aligned with the radio jet (e.g., de Vries \\etal~1999; Axon \\etal~2000) and that emission-line luminosities (e.g., \\oiii) are correlated with the radio power (e.g., Labiano 2008;  Buttiglione \\etal~2010; Kunert-Bajaraszewska \\& Labiano 2010). The optical emission-line diagnostics of the narrow-line region (NLR) can constrain accretion properties, since the NLR is photoionized by the nuclear continuum radiation (e.g., Kawakatu \\etal~2009).  In this work, using a sample of YRGs covering a large range of luminosities, we investigate the properties of NLRs and accretion by directly measuring narrow emission-line luminosities from optical spectra, and compare them with radio properties for constraining the disk-jet connection. In Section~\\ref{obs}, we describe the sample selection, spectroscopic observations, data reduction, and radio data collection. The measurements of emission-line fluxes and the stellar velocity dispersions are described in Section~\\ref{analysis}. Main results are presented in Section~\\ref{result} and Section~\\ref{sum} contains discussions and summary. Throughout the paper, we used cosmological parameters, $H_0 = 70$ \\kms~ Mpc$^{-1}$, $\\Omega_{m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$.  \\begin{deluxetable*}{llllllllrcccccr} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\tablecaption{Sample properties} \\tablehead{ \\colhead{Name} & \\colhead{R.A.} & \\colhead{Decl.} & \\colhead{z} & \\colhead{$E(B-V)$} & \\multicolumn{3}{c}{AGN type} & \\colhead{$S_{1.4}$}  & \\colhead{Ref.} & \\colhead{Jet size} & \\colhead{Ref.}  & \\colhead{Run} & \\colhead{Exposure} & \\colhead{S/N}\\\\ \\colhead{} & \\multicolumn{2}{c}{(J2000)} &  \\colhead{} &  \\colhead{}   &  \\colhead{} &  \\colhead{} &  \\colhead{} &  \\colhead{(Jy)} &  \\colhead{}&  \\colhead{(kpc)} &  \\colhead{} &  \\colhead{}  & \\colhead{(hr)} & \\colhead{} \\\\ \\colhead{(1)} & \\colhead{(2)} &  \\colhead{(3)} &  \\colhead{(4)} &  \\colhead{(5)} &  \\colhead{(6)} &  \\colhead{(7)} &  \\colhead{(8)} &  \\colhead{(9)} &  \\colhead{(10)} &  \\colhead{(11)} & \\colhead{(12)} & \\colhead{(13)} & \\colhead{(14)}  & \\colhead{(15)} } \\startdata \\multicolumn{15}{c}{Lick and Palomar targets}\\smallskip \\\\ \\hline 0019$-$000  \t& 00:22:25  & +00:14:56\t\t\t\t& 0.305\t\t&\t0.027\t\t\t\t&\t2\t\t  \t &\tGPS\t\t\t&\tHEG\t\t\t& 2.919 \t & F   \t& 0.121\t & O98\t\t\t&\t1 \t\t & 1.4\t\t\t& 48   \\\\ 0035+227 \t\t\t& 00:38:08  & +23:03:28\t\t\t\t& 0.096\t\t&\t0.033\t\t\t\t&\t2\t\t  \t &\tCSO\t\t\t&\tLEG\t\t\t& 0.547  \t & N   \t& 0.015\t & K08\t\t\t&\t3 \t\t & 2\t\t\t\t& 73   \\\\ 0134+329 \t\t\t& 01:37:41  & +33:09:35\t\t\t\t& 0.367\t\t&\t0.044\t\t\t\t&\t1\t  \t   &\tCSS\t  \t&\tHEG\t\t\t& 16.018\t & N   \t& 1.100  & O98\t\t\t&\t4 \t\t & 1\t\t\t\t& 208  \\\\ 0221+276 \t\t\t& 02:24:12  & +27:50:12\t\t\t\t& 0.310\t\t&\t0.125\t\t\t\t&\t1\t\t  \t &\tCSS\t\t\t&\tHEG\t\t\t& 3.024 \t & N   \t& 5.536  & K08\t\t\t&\t4 \t\t & 1.5\t\t\t& 32   \\\\ 0316+413 \t\t\t& 03:19:48  & +41:30:42\t\t\t\t& 0.018\t\t&\t0.163\t\t\t\t&\t1\t  \t   &\tCSO\t  \t&\tHEG\t\t\t& 22.829\t & N   \t& 0.005  & K08\t\t\t&\t3 \t\t & 1\t\t\t\t& 169  \\\\ 0345+337 \t\t\t& 03:48:47  & +33:53:15\t\t\t\t& 0.243\t\t&\t0.389\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t& 2.365 \t & N   \t& 0.443  & O98\t\t\t&\t5 \t\t & 3\t\t\t\t& 11   \\\\ 0428+205 \t\t\t& 04:31:04  & +20:37:34\t\t\t\t& 0.219\t\t&\t0.542\t\t\t\t&\t2\t\t  \t &\tGPS\t\t\t&\tLEG\t\t\t& 3.756 \t & N   \t& 0.412  & O98\t\t\t&\t2 \t\t & 1.5\t\t\t& 16   \\\\ 0605+480 \t\t\t& 06:09:33  & +48:04:15\t\t\t\t& 0.277\t\t&\t0.162\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t& 4.133 \t & N   \t& 15.17\t & K08\t\t\t&\t5 \t\t & 3\t\t\t\t& 12   \\\\ 0941$-$080\t\t& 09:43:37  & $-$08:19:31\t\t\t& 0.228\t\t&\t0.029\t\t\t\t& 2  \t\t\t &\tGPS\t\t\t&\tLEG\t\t  & 2.756 \t & N   \t& 0.083\t & O98\t\t\t&\t6 \t\t & 1.5\t\t\t& 10   \\\\ 1203+645 \t\t\t& 12:06:25  & +64:13:37\t\t\t\t& 0.372\t\t&\t0.017\t\t\t\t&\t1\t\t  \t &\tCSS\t\t\t&\tHEG\t\t\t& 3.719  \t & N   \t& 3.067\t & K08\t\t\t&\t10 \t\t & 3\t\t\t\t& 21   \\\\ 1225+442 \t\t\t& 12:27:42  & +44:00:42\t\t\t\t& 0.348\t\t&\t0.019\t\t\t\t&\t2\t\t  \t &\tGPS\t\t\t&\tHEG\t\t\t& 0.383 \t & F   \t& 0.493\t & K08\t\t\t&\t7 \t\t & 1.5\t\t\t& 9    \\\\ 1233+418 \t\t\t& 12:35:36  & +41:37:07\t\t\t\t& 0.250\t\t&\t0.022\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t& 0.664 \t & F   \t& 6.068\t & K10\t\t\t&\t7 \t\t & 0.5\t\t\t& 8    \\\\ 1245+676 \t\t\t& 12:47:33  & +67:23:16\t\t\t\t& 0.107\t\t&\t0.021\t\t\t\t&\t2\t\t  \t &\tCSO \t\t&\tLEG\t\t\t& 0.263 \t & N   \t& 0.007\t & K08\t\t\t&\t7 \t\t & 1.5\t\t\t& 141  \\\\ 1250+568 \t\t\t& 12:52:26  & +56:34:20\t\t\t\t& 0.320\t\t&\t0.010\t\t\t\t&\t1\t\t  \t &\tCSS\t\t\t&\tHEG\t\t\t& 2.442 \t & F   \t& 3.593\t & K08\t\t\t&\t8 \t\t & 1.5\t\t\t& 67   \\\\ 1323+321 \t\t\t& 13:26:17  & +31:54:10\t\t\t\t& 0.368\t\t&\t0.015\t\t\t\t&\t2\t\t  \t &\tGPS\t\t\t&\tHEG\t\t\t& 4.747 \t & F   \t& 0.149\t & K08\t\t\t&\t10 \t\t & 1.5\t\t\t& 34   \\\\ 1404+286 \t\t\t& 14:07:00  & +28:27:15\t\t\t\t& 0.077\t\t&\t0.018\t\t\t\t&\t1\t\t  \t &  GPS\t\t\t&\tHEG\t\t\t& 0.830 \t & F   \t& 0.005\t & K08\t\t\t&\t9 \t\t & 1\t\t\t\t& 53   \\\\ 1807+698 \t\t\t& 18:06:51  & +69:49:28\t\t\t\t& 0.051\t\t&\t0.036\t\t\t\t&\t2  \t\t\t &\tCSS \t\t&\tLEG\t\t\t& 1.886 \t & N   \t& 1.726\t & G94\t\t\t&\t3 \t\t & 1.5\t\t\t& 120  \\\\ 1943+546 \t\t\t& 19:44:32  & +54:48:07\t\t\t\t& 0.263\t\t&\t0.162\t\t\t\t&\t2\t\t  \t &\tCSO\t\t\t&\t\\nodata\t& 1.754 \t & N   \t& 0.072\t & K08\t\t\t&\t2 \t\t & 1.5\t\t\t& 13   \\\\ 2352+495 \t\t\t& 23:55:09  & +49:50:08\t\t\t\t& 0.238\t\t&\t0.181\t\t\t\t&\t2\t\t  \t &\tGPS\t\t\t&\tLEG\t\t\t& 2.306 \t & N   \t& 0.092\t & K08\t\t\t&\t2 \t\t & 2\t\t\t\t& 21 \\\\ \\cutinhead{SDSS targets} 0025+006 \t\t\t&\t00:28:33  & +00:55:11\t\t\t  & 0.104\t\t& 0.024\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tHEG\t\t\t&\t0.237 \t & F   \t& 2.230\t & K10\t\t\t&\t     \t & \\nodata\t& 21    \\\\ 0754+401 \t\t\t&\t07:57:57  & +39:59:36\t\t\t  & 0.066\t\t& 0.054\t\t\t\t&\t2\t\t  \t &  CSS\t\t\t&\tHEG\t\t\t&\t0.099 \t & F   \t& 0.178\t & K10\t\t\t&\t     \t & \\nodata\t& 28    \\\\ 0810+077 \t\t\t&\t08:13:24  & +07:34:06\t\t\t  & 0.112\t\t& 0.022\t\t\t\t&\t2  \t\t\t &\tCSS\t\t\t&\tLEG\t\t\t&\t0.463 \t & F   \t& 1.982\t & K10\t\t\t&\t     \t & \\nodata\t& 19    \\\\ 0921+143 \t\t\t&\t09:24:05  & +14:10:22\t\t\t  & 0.136\t\t& 0.029\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t&\t0.108    & F   \t& 0.520\t & K10\t\t\t&\t     \t & \\nodata\t& 19    \\\\ 0931+033 \t\t\t&\t09:34:31  & +03:05:45\t\t\t  & 0.225\t\t& 0.033\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t&\t0.292 \t & F   \t& 1.140\t & K10\t\t\t&\t     \t & \\nodata\t& 22    \\\\ 0942+355 \t\t\t&\t09:45:26  & +35:21:03\t\t\t  & 0.208\t\t& 0.011\t\t\t\t&\t1  \t\t\t &\tCSS\t\t\t&\tHEG\t\t\t&\t0.148  \t & F   \t& 3.138\t & K10\t\t\t&\t     \t & \\nodata\t& 20    \\\\ 1007+142 \t\t\t&\t10:09:56  & +14:01:54 \t\t\t& 0.213\t\t& 0.043\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t&\t1.045 \t & F   \t& 2.320\t & K10\t\t\t&\t     \t & \\nodata\t& 18    \\\\ 1037+302 \t\t\t&\t10:40:30  & +29:57:58\t\t\t  & 0.091\t\t& 0.019\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t&\t0.388  \t & F   \t& 2.591\t & K10\t\t\t&\t     \t & \\nodata\t& 30    \\\\ 1154+435 \t\t\t&\t11:57:28  & +43:18:06\t\t\t  & 0.230\t\t& 0.013\t\t\t\t&\t1\t\t  \t &\tCSS\t\t\t&\tHEG\t\t\t&\t0.256 \t & F   \t& 3.227\t & K10\t\t\t&\t     \t & \\nodata\t& 21    \\\\ 1345+125 \t\t\t&\t13:47:33  & +12:17:24\t\t\t  & 0.122\t\t& 0.034\t\t\t\t&\t1 \t\t\t &\tGPS\t\t\t&\tHEG\t\t\t&\t4.860  \t & F   \t& 0.085\t & O98\t\t\t&\t     \t & \\nodata\t& 17    \\\\ 1407+363 \t\t\t&\t14:09:42  & +36:04:16\t\t\t  & 0.148\t\t& 0.012\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tHEG\t\t\t&\t0.143 \t & F   \t& 0.050\t & K10\t\t\t&\t     \t & \\nodata\t& 9     \\\\ 1521+324 \t\t\t&\t15:23:49  & +32:13:50\t\t\t  & 0.110\t\t& 0.026\t\t\t\t&\t2  \t\t\t &\tCSS\t\t\t&\tHEG\t\t\t&\t0.169    & F   \t& 0.285\t & K10\t\t\t&\t     \t & \\nodata\t& 17    \\\\ 1558+536 \t\t\t&\t15:59:28  & +53:30:55\t\t\t  & 0.179\t\t& 0.012\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t&\t0.182    & F   \t& 3.630\t & K10\t\t\t&\t     \t & \\nodata\t& 17    \\\\ 1601+528 \t\t\t&\t16:02:46  & +52:43:58\t\t\t  & 0.106\t\t& 0.019\t\t\t\t&\t1\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t&\t0.576 \t & F   \t& 0.269\t & K10\t\t\t&\t     \t & \\nodata\t& 26    \\\\ 1610+407 \t\t\t&\t16:11:49  & +40:40:21 \t\t\t& 0.151\t\t& 0.008\t\t\t\t&\t2\t\t  \t &\tCSS\t\t\t&\tLEG\t\t\t&\t0.553 \t & F   \t& 1.858\t & K10\t\t\t&\t     \t & \\nodata\t& 12 \\enddata \\tablecomments{Columns: (1) Target name; (2) R.A.; (3) Decl..; (4) Redshift; (5) Galactic extinction; (6) Spectroscopic AGN type -- 1: Type 1 AGN with broad emission lines; 2: Type 2 AGN without broad emission line; (7) YRG type; (8) Classification by excitation -- HEG: high excitation galaxy, LEG: low excitation galaxy; (9) 1.4 GHz integrated flux density of radio source; (10) Reference for the flux density -- F: from the Faint Images of the Radio Sky at Twenty-centimeters (FIRST) catalog, N: from the NRAO/VLA Sky Survey (NVSS) catalog; (11) Jet size; (12) Reference for the jet size -- G94: Gelderman \\& Whittle (1994), O98: O'Dea (1998), K08: Kawakatu \\etal~(2008); K10: Kunert-Bajaraszewska \\etal~(2010); (13) Observing run in Table~\\ref{setup};  (14) Exposure time; (15) Signal-to-noise ratios near 5100 \\AA\\ in the rest frames.} \\label{targets} \\end{deluxetable*} ",
        "conclusions": "{\\label{sum}}  To investigate spectral properties of YRGs and compare them with radio properties, we construct a sample of 34 YRGs at relatively low redshift ($z < 0.4$) for measuring narrow emission-line properties, \\mbh, and Eddington ratio. We determined \\mbh\\ from the width and luminosity of the broad \\Ha\\ line using single-epoch mass estimators for Type 1 (broad-line) AGNs, or from the measured stellar velocity dispersion using the \\msigma\\ relation for Type 2 (narrow-line) AGNs. The estimated \\mbh\\ ranges from 10$^{7.0}$ to 10$^{9.2}$ \\msun, indicating YRGs have relatively massive BHs, similar to the large-scale radio galaxies.  Based on the narrow emission-line flux ratios (e.g. \\oiii/\\Hb, [N {\\sc ii}]/\\Ha, [S {\\sc ii}]/\\Ha, and \\oi/\\Ha), we classified YRGs as HEG and LEG. Most of Type 1 AGNs belong to HEGs, while Type 2 AGNs are composed of HEGs and LEGs. We find that the Eddington ratio of HEGs is higher by $\\sim$1.0 dex than that of LEGs, using the \\Ha\\ line luminosity as a proxy for AGN bolometric luminosity. The difference in Eddington ratios and comparison with photoionization models suggest that HEGs are high Eddington ratio AGNs with an optically thick accretion disk, which are similar to QSOs or Seyfert 1 galaxies, while LEGs have lower Eddington ratios with radiatively inefficient accretion flow. This interpretation is similar to the division between Seyfert galaxies and LINERs in RQ AGNs (Kewley \\etal~2006; Ho 2008), suggesting that YRGs have a various range of accretion activities over 2-3 orders of magnitude in the Eddington ratio.  Kawakatu \\etal~(2009) investigated whether the optical narrow emission-line ratios of YRGs are systematically different from those of RQ Seyfert 2 galaxies by comparing the observed line ratios (e.g., [O {\\sc i}]/[O {\\sc iii}] and [O {\\sc ii}]/[O {\\sc iii}]) with photoionization models. Using a limited sample of YRGs, they concluded that YRGs favor SED without a strong BBB, i.e., optically thin advection-dominated accretion flow, while RQ AGNs are consistent with the models adopting SED with a strong BBB, i.e., a geometrically thin, optically thick disk. In this study with an enlarged sample including Type 1 AGNs with higher Eddington ratios, we find that there are various levels of accretion activity in YRGs and that both SEDs with/without BBB are required to reproduce the observed line flux ratios of YRGs.  Low luminosity AGNs, i.e., LINERs, generally tend to be radio-loud (Ho 1999, Terashima \\& Wilson 2003), implying that radio activity may be related with radiatively inefficient accretion flow, similar to the low-state in X-ray binaries (McClintock \\& Remillard 2006). In the case of YRGs, we find a large range of Eddington ratios, including HEGs with high accretion power and Seyfert-like emission-line flux ratios. Thus, the connection between radio jet and radiatively inefficient accretion flow is not strong in YRGs. Instead, YRGs are probably composed of heterogeneous objects representing various accretion states.  By comparing narrow emission-line properties with radio luminosity and jet size, we investigated the disk-jet connection in YRGs. The \\oiii\\ and \\Ha\\ line luminosities show broad correlations with the radio core luminosity, indicating that accretion and radio activities in YRGs are fundamentally linked. However, at fixed radio luminosity, HEGs have higher line luminosities (particularly for \\oiii) than LEGs, indicating that HEGs have higher accretion activity than LEGs for a given radio activity. These results may suggest that at a given radio activity there is a continuous distribution of accretion powers due to various mass accretion rate."
    },
    "1208/1208.6576_arXiv.txt": {
        "abstract": "We report on radio and X-ray observations of PSR~\\psr, a 533-ms radio pulsar discovered in the Parkes Multibeam Pulsar Survey. From radio observations taken with the Parkes, Lovell and Arecibo telescopes, we show that this pulsar exhibits two spindown states akin to PSRs B1931+24 reported by Kramer et al. and J1841$-$0500 reported by Camilo et al.  Unlike PSR~B1931+24, which switches between ``on'' and ``off'' states on a 30--40 day time-scale, PSR~\\psr\\ is similar to PSR~J1841$-$0500 in that it spends a much longer period of time in the off-state. So far, we have fully sampled two off-states. The first one lasted between 560 and 640 days and the second one lasted between 810 and 835 days. From our radio timing observations, the ratio of on/off spindown rates is $1.77 \\pm 0.03$.  {\\it Chandra} observations carried out during both the on- and off-states of this pulsar failed to detect any emission. Our results challenge but do not rule out models involving accretion onto the neutron star from a low-mass stellar companion. In spite of the small number of intermittent pulsars currently known, difficulties in discovering them and in quantifying their behavior imply that their total population could be substantial. ",
        "introduction": "\\label{sec:intro}\\setcounter{footnote}{0}  It is well established that not all radio pulsars emit radiation during each rotation. Backer (1970) \\nocite{bac70} first observed this phenomenon and demonstrated that some pulsars exist in a ``null state'' for several pulse periods before switching back on again. Pulsar nulling has been investigated extensively over the years (e.g., Ritchings 1976; Rankin 1986; Biggs 1992). \\nocite{rit76,ran86,big92a} From a study of pulsars in the Parkes Multibeam Pulsar Survey (PMPS; Manchester \\nocite{mlc+01} et al.~2001), \\nocite{wmj07} Wang et al.~(2007) confirmed earlier evidence (Ritchings 1976) that the fraction of nulling pulses generally increases with increasing characteristic age.  In addition to the nulling phenomenon, it has become apparent that a new class of intermittent pulsars exist where no radiation is observed over much longer time-scales.  In a single-pulse analysis of archival PMPS data, McLaughlin et al.~(2006) \\nocite{mll+06} discovered a new class of neutron stars (Rotating Radio Transients) from which radio emission is detectable, on average, only 1~s per day in an apparently random fashion. In the same year, Kramer et al.~(2006)\\nocite{klo+06} reported the discovery of a more deterministic type of intermittency in PSR~B1931+24, which appears to be the prototype of a large population of pulsars that have so far been difficult to detect. As Kramer et al.~demonstrated, PSR~B1931+24 shows a quasi-periodic on/off cycle with a period of 30--40 days in which the spindown rate increases by $\\sim 50$\\% when the pulsar is in its on-state compared to the off-state. In this paper, we report observations characterizing intermittent behavior in PSR~\\psr\\ an apparently ordinary 533-ms pulsar with a characteristic age of 5.6~Myr which was discovered as part of the PMPS (Lorimer et al.~2006). \\nocite{lfl+06} Earlier accounts of this work we presented by Kramer (2008) and Lyne (2009).  \\nocite{kra08,lyn09} Very recently, Camilo et al.~(2012) \\nocite{crc+12} announced the discovery of PSR~J1841$-$0500, a 912-ms pulsar which has so far shown one off-state lasting 580~days. Like B1931+24, the spindown rate in the on-state is higher than the off-state. For J1841$-$0500 the increase is approximately 150\\%! These pulsars are dramatic examples of a newly recognized and large group of pulsars which show changes in their emission properties and period derivatives \\citep{lhk+10} which are correlated and often quasi-periodic. Understanding pulsar intermittency will shed new insights into neutron star physics and populations.  The long off-states of intermittent pulsars are in stark contrast to the longest known quiescence times of nulling pulsars, i.e.~they exceed the typical nulling time scale by about five orders of magnitude. In addition, the observed increase in spindown rate points to a significant increase in the magnetospheric particle outflow when the pulsar switches on, indicating that a pulsar wind plays a significant role in neutron star spin evolution. As described by Kramer et al.~(2006), and discussed later in this paper, the spindown rate changes allow us to estimate the current density associated with the radio emission.  The difficulties in detecting and identifying intermittent pulsars imply that the few we currently observe represent a potentially substantial population of similar objects in the Galaxy. To better understand this population, it is therefore important to establish the related time-scales for the non-emitting state.  Here we detail our observations of intermittent behavior in PSR~\\psr. In \\S \\ref{sec:Radio} we describe the radio observations we have carried out to characterize its intermittency. In \\S \\ref{sec:Xray} we describe the {\\it Chandra} X-ray observations which constrain the high-energy emission from the pulsar. In \\S \\ref{sec:discussion}, we discuss the implications of our results. ",
        "conclusions": "\\label{sec:discussion}  The radio observations reported above suggest long-term intermittent behavior in PSR~\\psr, while the X-ray observations failed to detect any difference in the high-energy emission between the on- and off-states. We now discuss the implications of these observations.  \\subsection{Spindown behavior in the two states}  To track the variation in spin frequency ($\\nu$) of PSR~\\psr\\ we used the {\\sc tempo2} software package (Hobbs et al.~2006) \\nocite{hem06} and its {\\tt stridefit} plugin to carry out measurements of $\\nu$ based on timing model fits to short (30-day time span) segments of data in which the position was fixed at the nominal values found by Lorimer et al.~(2006) and the spin frequency derivative $\\dot{\\nu}$ was assumed to be zero. As shown in Fig.~\\ref{fig:spindown}, this analysis reveals a clear discontinuity in the spindown behavior during the off-states. Two explanations could account for this: (i) the pulsar suffered a period glitch during the off-states; (ii) the spindown rate was different during the off-states, as is observed for PSRs~B1931+24 and J1841$-$0500. Although it is possible to fit across the 2004--2005 period, resulting in $\\Delta \\nu/\\nu=(5.34 \\pm 0.07) \\times 10^{-8}$ for the putative glitch, no exponential recovery is observed and the abrupt turn-off observed in emission is inconsistent with other observations of glitching pulsars \\citep[see, e.g.,][]{elsk11}. A similar conclusion can be reached for the 2010--2012 off-state. Henceforth we examine the hypothesis where the spindown rate of PSR~\\psr\\ has two distinct values which we refer to as $\\dot{\\nu}_{\\rm on}$ and $\\dot{\\nu}_{\\rm off}$.  Using {\\sc tempo2}, we obtained independent timing solutions in each of the two on-states. The results of these analyses are summarized in Table 1. The pulsar's DM was not constrained by these analyses and was therefore held fixed at the value reported by Lorimer et al.~(2006; DM = 28.3~cm$^{-3}$~pc).  The shorter timing baseline ($\\sim 270$~days) for the first on-state compared to the second one ($\\sim 4$~yr) means that the timing parameters obtained from it are less precise and subject to covariances. So far, we have only sampled $\\sim 1$~month in the current (third) on-state. To minimize these covariances, we held the position in the first on-state fit fixed at the position derived from the second on-state. While the post-fit residuals for the first on-state are approximately white, a significant amount of timing noise is present in the residuals from the second on-state shown in Fig.~\\ref{fig:residuals}. This behavior can be removed by fitting multiple sinusoids to the data (the ``harmonic whitening'' technique; \\nocite{hlk+04} Hobbs et al.~2004). To check the effect on the measured parameters in Table 1, we carried out such an analysis using the {\\tt fitwaves} plugin to {\\sc tempo2}. The residuals can be whitened by removing six harmonically related sinusoids and the result fit parameters are all within 1-$\\sigma$ of the values presented in Table 1. We therefore adopt the parameters from the second on-state as being our most precise measurements of the pulsar to date and, hence, $\\dot{\\nu}_{\\rm on}=-(5.44505 \\pm 0.00007) \\times 10^{-15}$~s$^{-2}$.  To measure $\\dot{\\nu}_{\\rm off}$, we accounted for the uncertainties in off/on switching epochs in the following way. We first assumed that the nominal switch-off epoch ($T_{\\rm off}$) of the pulsar occured midway between the date of the last detection during the second on-phase ($T_1 =$~MJD 55293.33) and the first non-detection ($T_2 =$~MJD 53301.99). Similarly, for the nominal switch-on epoch ($T_{\\rm on}$), we adopt the midpoint between the last non-detection ($T_3=$~MJD 56110.85) and the first re-detection of the third on-phase ($T_4=$~MJD 56128.78). Using {\\sc tempo2}, we computed $\\nu(T_{\\rm off})$ and $\\nu(T_{\\rm on})$ --- the nominal pulse frequencies at both $T_{\\rm off}$ and $T_{\\rm on}$ as predicted by the second and third on-state timing models, respectively.  The off-state spindown rate \\begin{displaymath} \\dot{\\nu}_{\\rm off} = \\frac{\\nu(T_{\\rm on})- \\nu(T_{\\rm off})}{T_{\\rm on}-T_{\\rm off}} = -(3.08 \\pm 0.05) \\times 10^{-15}\\,{\\rm s}^{-2}. \\end{displaymath} Here the uncertainty in $\\dot{\\nu}_{\\rm off}$ is dominated by the uncertainty in $T_{\\rm on}-T_{\\rm off}$ which we estimated to be the mean of the two time windows of interest here (i.e. $(T_4-T_3+T_2-T_1)/2)$). A similar analysis for the first off-state yields $\\dot{\\nu}_{\\rm off} = -(3.2 \\pm 0.2) \\times 10^{-15}$~s$^{-2}$. These results imply that the ratio of on/off spindown rates ${\\cal R}=\\dot{\\nu}_{\\rm on}/\\dot{\\nu}_{\\rm off}$ is therefore $1.77 \\pm 0.03$, i.e.~slightly higher than PSR~B1931+24 but below PSR~J1841$-$0500.  \\subsection{Charge density in the pulsar magnetosphere} \\label{sec:magnetosphere}  To estimate the implied current flow in the pulsar magnetosphere for both the on and off-states, we follow Kramer et al.~(2006) and consider the simplest possible emission model. We assume that, in the off-state, the pulsar spins down by a mechanism that does not involve a substantial particle ejection (e.g., it would be magnetic dipole radiation if the pulsar were in vacuum), while the rate in the on-state is enhanced by a torque from the current of an additional plasma outflow.  Assuming that the spindown energy loss rate in the on-state, $\\dot{E}_{\\rm on}$, may be written as the sum of the energy loss rate in the off-state, $\\dot{E}_{\\rm off}$, and the energy loss due to the additional plasma, $\\dot{E}_{\\rm plasma}$, the corresponding charge density (in cgs units) in the plasma \\begin{equation} \\rho_{\\rm plasma} = \\frac{3 I (\\dot{\\nu}_{\\rm off} - \\dot{\\nu}_{\\rm on})}{R^4_{\\rm pc} B_{\\rm off}}. \\end{equation} Here $I$ is the moment of inertia of the neutron star, \\begin{equation} R_{\\rm pc} = \\sqrt{\\frac{2 \\pi \\nu R^3_{\\rm NS}}{c}} \\end{equation} is the polar cap radius, $B_{\\rm off}$ is the dipole surface magnetic field strength calculated from the spin frequency and spindown rate in the off-state, and $c$ is the speed of light. For a canonical neutron star of radius $R_{\\rm NS}=10^6$~cm and moment of inertia $I=10^{45}$~gm~cm$^2$, we find $\\rho_{\\rm plasma} \\simeq 62$~esu~cm$^{-3}$.  This is slightly higher than the so-called Goldreich-Julian density \\nocite{gj69} \\begin{equation} \\rho_{\\rm GJ} = \\frac{B_{\\rm off}}{Pc} \\simeq 44 \\,{\\rm esu}\\,{\\rm cm}^{-3} \\end{equation} which is the charge density required to radiate along the open magnetic field lines in the idealized pulsar magnetosphere model proposed by Goldreich \\& Julian (1969).  For PSR~B1931+24, Kramer et al.~(2006) found that $\\rho_{\\rm GJ} \\simeq \\rho_{\\rm plasma} \\simeq 100$~esu~cm$^{-3}$. Taking the corresponding values for PSR~J1841$-$0500 from Camilo et al.~(2012), we find for this pulsar that the plasma density $\\rho_{\\rm plasma} \\simeq 400$~esu~cm$^{-3}$ is significantly larger than $\\rho_{\\rm GJ} \\simeq 130$~esu~cm$^{-3}$. The fact that these inferred densities all equal or exceed $\\rho_{\\rm GJ}$ at least implies that the basic conditions for radiation by the Goldreich \\& Julian (1969) model are being met.  For simplicity, the above calculations make the assumption that the pulsar is an orthogonal rotator. In reality, of course, the inclination angle between the spin and magnetic axes $\\alpha < 90^{\\circ}$.  More recent and realistic modeling of the pulsar magnetosphere by Li et al.~(2012) and Kalapotharakos et al.~(2012) \\nocite{lst12,kkhc12} consider force-free electrodynamic and resistive solutions which can account for the different spin-down rates observed in these three intermittent pulsars. Based on our measurement of the on/off spindown ratio, ${\\cal R}$, and the results presented in Fig.~3 of Li et al.~(2012), the prediction for PSR~\\psr\\ is that $\\alpha \\sim 60^{\\circ}$. A future Arecibo observing campaign on this pulsar during its next on-state will be undertaken to obtain high-quality polarimetric data with the aim of constraining $\\alpha$.  As discussed by Beskin \\& Nokhrina (2007) \\nocite{bn07} and Kalapotharakos et al.~(2012), further discoveries of intermittent pulsars with different values of ${\\cal R}$ than observed so far would greatly help to constrain these models.  In the context of models for pulsar intermittency involving the neutron star's emission mechanism, it should also be noted that Zhang, Gil \\& Dyks~(2007) \\nocite{zgd07} proposed that intermittent pulsars are old isolated neutron stars which have entered the so-called ``death valley'' in the $P-\\dot{P}$ diagram (Chen \\& Ruderman 1993) \\nocite{cr93} where the voltage across the neutron star polar cap is no longer sufficient for pair production in the neutron star magnetosphere, and the radio emission becomes sporadic. Zhang et al.\\ suggest that non-dipolar magnetic field configurations, similar to the sunspot phenomenon, may be effective in such neutron stars and temporarily rejuvenate their radio emission. It is not clear, however, how the quasi-periodic nature seen in PSR~B1931+24 can be explained quantitatively in this scenario, or indeed whether it applies to PSR~\\psr\\ or PSR~J1841$-$0500, since \\citep[as noted by][]{crc+12} none of these pulsars are in the death valley region.  \\subsection{Constraints from the X-ray non-detections} \\label{sec:xraycons}  The X-ray count rate upper limits can be used to estimate upper limits on energy flux, which, however, depend on assumed spectrum. We know from observations of old pulsars that their X-ray spectra can be approximated by an absorbed power-law model with a photon index $\\Gamma \\approx 2$--4 (e.g., Kargaltsev et al.\\ 2006; Pavlov et al.\\ 2009). \\nocite{kpg06,pkwg09} The hydrogen column density toward the pulsar, $N_{\\rm H} \\approx 1 \\times 10^{21}$ cm$^{-2}$, can be estimated from the DM assuming a 10\\% average degree of ionization of the interstellar medium. Using the {\\sl Chandra} PIMMS tool\\footnote{http://asc.harvard.edu/toolkit/pimms.jsp}, we obtain the following absorbed and unabsorbed energy fluxes for a given source count rate $R_s$ in the first observation: $F_{0.3-8\\,{\\rm keV}}^{\\rm abs} = 0.63$, 0.47 and 0.48, $F_{0.3-8\\,{\\rm keV}}^{\\rm unabs} = 0.82$, 0.90 and 1.33, in units of $10^{-15} (R_{s}/10^{-4}\\,{\\rm counts}~{\\rm s}^{-1})$~erg~cm$^{-2}$ s$^{-1}$, for $\\Gamma =2$, 3 and 4, respectively. For the second observation, the corresponding fluxes are $F_{0.3-8\\,{\\rm keV}}^{\\rm abs} = 0.70$, 0.55 and 0.59, $F_{0.3-8\\,{\\rm keV}}^{\\rm unabs} = 0.91$, 1.04 and 1.64, in the same units. Note that the same count rates correspond to higher fluxes in the second observation because the ACIS effective area became smaller.  Using these relations and the count rate upper limits estimated above, we can estimate the flux upper limits at a given confidence level.  For instance, for $\\Gamma=3$ and $C=0.99$, we obtain $F_{0.3-8\\,{\\rm keV}}^{\\rm abs} < 0.7$ and $<2.0$, $F_{0.3-8\\,{\\rm keV}}^{\\rm unabs} < 1.3$ and $<3.8$, for the first and second observations, respectively, in units of $10^{-15}$ erg cm$^{-2}$ s$^{-1}$.  From these upper limits, one can estimate upper limits on X-ray luminosity, $L_X = 4\\pi d^2 F_X^{\\rm unabs}$ and efficiency, $\\eta_X = L_X/\\dot{E}$. For the on-state (where $\\dot{E} = 4.0\\times 10^{32}$ erg s$^{-1}$), assuming $\\Gamma=3$, we obtain $L_{\\rm 0.3-8\\,keV} < 2.6\\times 10^{29} (d/1.3\\,{\\rm kpc})^2$ erg~s$^{-1}$, $\\eta_{\\rm 0.3-8\\,keV} < 6.5\\times 10^{-4} (d/1.3\\,{\\rm kpc})^2$, at $C=0.99$. Making the same assumptions for the off-state (where $\\dot{E} = 2.4\\times 10^{32}$ erg s$^{-1}$), we obtain $L_{\\rm 0.3-8\\,keV} < 7.7\\times 10^{29} (d/1.3\\,{\\rm kpc})^2$ erg~s$^{-1}$, $\\eta_{\\rm 0.3-8\\,keV} < 3.0\\times 10^{-3} (d/1.3\\,{\\rm kpc})^2$, at $C=0.99$. These limits are consistent with the non-thermal efficiencies observer in other non-recycled pulsars \\citep{pccm02,zp04}.  \\subsection{Implications for other intermittency models}  Our discussion so far has focused on pulsar intermittency as being due to processes that are internal to the neutron star magnetosphere.  At least two alternative scenarios, which we discuss below, have been made to explain the phenomenon as being due to the influence of material emanating from outside the magnetosphere.  Cordes \\& Shannon (2008) \\nocite{cs08} investigated the consequences of debris disks around neutron stars, i.e.~metal-rich leftover material from the supernova explosion that has aggregated into a disk of circumpulsar material. They propose a scenario in which the behavior seen in PSR~B1931+24 is produced by an asteroid in an eccentric 40-day orbit which deflects material from the debris disk into the neutron star magnetosphere. Such a process could temporarily halt the electron-positron pair production thought to be responsible for the radio emission.  Unfortunately, the infall rates required by this model translate to completely undetectable X-ray fluxes.  In addition, given that sufficiently high-precision timing is not possible for pulsars such as PSR~\\psr, any periodic signatures from such small bodies would not be detectable in its timing residuals (Fig.~\\ref{fig:residuals}).  Rea et al.\\ (2008) \\nocite{rks+08} suggested that accretion onto the neutron star from a low-mass stellar companion in an eccentric orbit close to periastron could halt pair production. In this case, the heating of the infalling matter would produce additional X-rays in the off-state. While no signatures indicative of a binary companion exist in our radio timing residuals, such an orbit would not be detectable if it were close to face-on. Rea et al.\\ attempted to test this hypothesis for PSR~B1931+24 via a {\\it Chandra} ACIS observation in 2006. Unfortunately, the pulsar switched on unexpectedly before their observations.  To test this model in our observations of PSR~\\psr, following the discussion in \\S 5.2 from Rea et al.\\ (2008), we assume that the radio emission is quenched when the neutron star's Alfven radius is less than its light cylinder radius. This corresponds to $L_X \\gapp 10^{30}$~erg~s$^{-1}$ and is right at the boundary of detectability in our off-state observation given the upper limit $L_{\\rm 0.3-8\\,keV} < 7.7\\times 10^{29} (d/1.3\\,{\\rm kpc})^2$ erg~s$^{-1}$ found in \\S \\ref{sec:xraycons}. However, since the distance estimate to PSR~\\psr\\ made using the Cordes \\& Lazio (2002) electron density model can be uncertain by factors of two or more \\citep[see, e.g.,][]{dtbr09} we cannot therefore conclusively reject this scenario as an explanation for the behavior observed in PSR~\\psr.  \\subsection{How common are intermittent pulsars?}  Regardless of the form of the mechanism for pulsar intermittency, its recognition and characterization through the three pulsars so far poses interesting questions as to the size of the likely population of similar objects in the Galaxy. From our sampling of PSR~\\psr\\ so far, it appears to spend approximately 50\\% of the time in the off-state. Similar considerations for PSRs~B1931+24 and J1841$-$0500 imply similar off-state duty cycles. Due to these and similar pulsars being less likely to be on during pulsar search and confirmation observations, as noted by Kramer et al.~(2006), they could represent a substantial population that has so far evaded detection. PSR~J1841$-$0500, for example, was in an off-state during the closest PMPS observation, and was only discovered serendipitously during a search of a magnetar in the same telescope beam \\citep{crc+12}. In addition to intermittent pulsars evading discovery in large-scale surveys, since a significant fraction of the $\\sim$ 1700 non-recycled pulsars currently known are not subject to long-term timing programs, it is currently unclear as to what fraction of these could exhibit intermittency on long time scales. Even the prototypical object B1931+24 evaded characterization for almost 20 years after its discovery \\nocite{stwd85} (Stokes et al.~1985). For PSR~\\psr, had the initial Parkes timing observations spanned a period of a year and no off-state seen, it is possible that the intermittent behavior would have evaded detection.  Perhaps the majority of normal pulsars exhibit some form of intermittent behavior, if they are studied long enough.  If that is the case, then current estimates of the pulsar birthrate would need to be upwardly revised by a factor of two. The impact of the discovery of intermittent pulsars on our understanding of the neutron star birth rate \\citep[see, e.g.,][]{kk08} is a currently unsolved problem which merits further investigation.  \\bigskip The Parkes radio telescope is part of the Australia Telescope, which is funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO.  The Arecibo Observatory is operated by SRI International under a cooperative agreement with the National Science Foundation (AST-1100968), and in alliance with the Ana G. M\\'endez-Universidad Metropolitana, and the Universities Space Research Association.  We thank Arun Venkataraman for retrieving the Arecibo observations.  D.R.L., G.G.P., C.C., and M.A.M. were supported by {\\it Chandra} grants GO8-9078 and GO1-12078 during part of this work. D.R.L. and M.A.M. were also supported by a WVEPSCoR Research Challenge Grant and acknowledge the support provided by the Australia Telescope National Facility distinguished visitor program during completion of this work. The work of G.G.P. was also partly supported by NASA grant NNX09AC84G and by the Ministry of Education and Science of Russian Federation (contract 11.G34.31.0001). We thank Fernando Camilo and Andrew Seymour for useful comments on the manuscript, as well as Christine Jordan for assistance with the Lovell timing data and George Hobbs for useful discussions concerning {\\sc tempo2}."
    },
    "1208/1208.3679_arXiv.txt": {
        "abstract": "Early in the reionization process, the intergalactic medium (IGM) would have been quite inhomogeneous on small scales, due to the low Jeans mass in the neutral IGM and the hierarchical growth of structure in a cold dark matter Universe. This small-scale structure acted as an important sink during the epoch of reionization, impeding the progress of the ionization fronts that swept out from the first sources of ionizing radiation. Here we present results of high-resolution cosmological hydrodynamics simulations that resolve the cosmological Jeans mass of the neutral IGM in representative volumes several Mpc across. The adiabatic hydrodynamics we follow are appropriate in an {\\em unheated} IGM, before the gas has had a chance to respond to the photoionization heating. Our focus is determination of the resolution required in cosmological simulations in order to sufficiently sample and resolve small-scale structure regulating the opacity of an unheated IGM. We find that a dark matter particle mass of $m_{\\rm dm} \\lesssim 50~M_\\odot$ and box size of $L \\gtrsim 1$ Mpc are required. With our converged results we show how the mean free path of ionizing radiation and clumping factor of ionized hydrogen depends upon the ultraviolet background (UVB) flux and redshift. We find, for example at $z=10$, clumping factors typically of 10 to 20 for an ionization rate of $\\Gamma \\sim 0.3 - 3 \\times 10^{-12}$s$^{-1}$, with corresponding mean free paths of $\\sim 3-15$~Mpc, extending previous work on the evolving mean free path to considerably smaller scales and earlier times. ",
        "introduction": "The fact that the most abundant sources of radiation during reionization are likely to be currently undetectable \\citep[e.g.,][]{trenti/etal:2010, oesch/etal:2012, alvarez/etal:2012} means that the details of the reionization process are beyond most current observational probes. The notable exceptions are observations of the polarization of the cosmic microwave background (CMB), which imply an optical depth to Thomson scattering of $\\tau\\sim 0.09$ \\citep{komatsu/etal:2011}, and the appearance of a Gunn-Peterson trough in the spectra of distant quasars \\citep{Fan06}, indicating that reionization was largely complete by $z\\sim 6$. Reionization is therefore thought to have mainly taken place over the redshift range $z\\sim 6-15$. Due to the lack of more specific constraints, much of our current understanding about the epoch of reionization comes from theoretical studies in the context of the $\\Lambda$CDM cosmology.  The picture which always emerges is of small-scale gaseous structures forming at $z>20$, due to the collapse of dark matter halos at the Jeans scale, roughly $10^4~M_\\odot$ \\citep[e.g.,][]{peebles/dicke:1968,couchman/rees:1986,shapiro/etal:1994,gnedin/hui:1998}. The gas was just cool enough to fall into halos at this mass, leading to strong inhomogenities on a scale of tens of comoving parsecs. At the same time, slightly more massive halos, with masses on the order of $\\sim 10^6~M_\\odot$, formed enough H$_2$ molecules in their cores to cool efficiently, leading to the formation of the first stars in the Universe \\citep[e.g.,][]{tegmark/etal:1997,2002Sci...295...93A,2002ApJ...564...23B,2003ApJ...592..645Y}. The ionizing radiation from these stars is thought to have created substantial, yet short-lived H~II regions, which were shaped by the surrounding inhomogeneity of the  gas distribution \\citep{alvarez/etal:2006,abel/etal:2007,yoshida/etal:2007}.  Eventually, sufficiently large halos formed that triggered the formation of the first galaxies \\citep{johnson/etal:2007,wise/abel:2008,greif/etal:2008}. These nascent dwarf galaxies would have created longer-lived and isolated H~II regions \\citep{wise/cen:2009,wise/etal:2012}. It is unclear how these galaxies evolved into the much more luminous ones that have been observed at redshifts as high as $z\\sim 8$ \\citep[e.g.,][]{bouwens/etal:2010}.  Nevertheless, it is widely believed that as the first galaxies grew and merged, their collective radiative output created a large and complex patchwork of ionized bubbles, with characteristic sizes on the order of tens to hundreds of comoving Mpc \\citep[e.g.,][]{barkana/loeb:2004,furlanetto/etal:2004,iliev/etal:2006}. During this time, dense systems in the IGM likely impeded the progress of ionization fronts \\citep{barkana/loeb:1999,haiman/etal:2001,shapiro/etal:2004,iliev/etal:2005,ciardi/etal:2006}. At the end of reionization the so-called ``Lyman-limit'' systems, dense clouds of gas optically-thick to ionizing radiation observed in the spectra of quasars at $z<6$ \\citep[e.g.,][]{storrie-lombardi/etal:1994,prochaska/etal:2009}, dominated the overall opacity of the IGM to ionizing radiation. These systems crucially influenced the percolation phase of reionization \\citep{gnedin/fan:2006,choudhury/etal:2009,alvarez/abel:2012}, which in turn determined the evolution and structure of the ionizing background \\citep[e.g.,][]{haardt/madau:1996,bolton/haehnelt:2007a,mcquinn/etal:2011}.  Thus, the progress of reionization depended not only on the properties of the sources of ionizing radiation, but also on the {\\em sinks}. Theoretical models of reionization must describe not just the spectral energy distribution, abundance, and clustering of early sources of ionizing radiation, but also the inhomogeneity of the intergalactic medium (IGM) in the space between the sources. It is this latter description that is the goal of the present work.  Early descriptions of reionization took into account inhomogeneities in the IGM through a ``clumping factor'', $c_l$, by which the recombination rate is boosted relative to the homogeneous case. This allows one to write a global ionization rate, equal to the ionizing photon emmissivity minus the recombination rate of a clumpy IGM, and thereby determine the reionization history for a given ionizing source population. \\citet{shapiro/giroux:1987} used such a model to show that the observed population of QSOs were insufficient to have reionized the Universe by $z\\sim 5$. Their assumption of  $c_l\\sim 1$ would have been conservative, in that that additional recombinations would have made it even more difficult for quasars alone to reionize the Universe.  In addition to being useful in modeling the reionization history, the clumping factor is also important in estimating the necessary number of ionizing photons per baryon to {\\em maintain} an ionized Universe. The necessary and sufficient condition for maintaining an ionized Universe is that the ionizing photon emissivity should be greater than or equal to the recombination rate of the IGM. \\citet{madau/etal:1999} used this fact to derive a critical star formation rate, above which the rate of ionizing photons is enough to maintain the Universe in an ionized state.  \\citet{gnedin/ostriker:1997} used hydrodynamic simulations with a treatment of photoionization in the ``local optical depth'' approximation to determine the clumping factor of the ionized component of the IGM, finding a value of $c_l\\sim 30$ at $z=6$. They also pointed out that the actual clumping factor of the IGM would have been larger due to structure on smaller scales than they resolved. More recently, \\citet{miralda-escude/etal:2000} built a semi-analytical model for the reionization of an inhomogenous IGM, in which the underlying gas density distribution was determined by numerical simulations.  They argued that in addition to specifying the clumping factor of the ionized medium, it is also necessary to describe the distribution of high-density gas clouds that are able to self-shield against ionizing radiation.  \\cite{mcquinn/etal:2011} followed a similar approach to that of \\cite{miralda-escude/etal:2000} to  explain the evolution of the ionizing background radiation at redshifts less than $z\\sim 6$, using more realistic numerical simulations which were post-processed with radiative transfer. These works were focused on the large scales relevant in the post-reionization IGM, after photoionization heating has ``ironed out'' the clumpiness of the IGM on the smallest scales. The timescale over which this smoothing occurs is on the order of $10-100$ Myr \\citep{2005MNRAS.361..405I}. Although the recombination rate in the homogenous IGM is on the order of 1 Gyr, small-scale inhomogeneities increase recombinations by at least an order of magnitude, making the recombination time of the high-redshift IGM comparable to the smoothing time. Our work here is focused on the higher redshifts and smaller scales that were most relevant early in reionization before much smoothing has occurred.  We seek to obtain convergence in the quantities that describe the inhomogeneity of the {\\em unheated} IGM during the epoch of reionization, such as the mean free path, $\\lambda$, clumping factor, $c_l$, and density threshold above which gas is self-shielded, $\\Delta_{\\rm crit}$, by spanning the parametre space of redshift and ionizing background intensity, $j_\\nu$. To do this, we post-process cosmological adiabatic hydrodynamics simulations with radiative transfer calculations along different lines of sight through the simulated volume. Radiative feedback raises the Jeans mass of the IGM, thereby increasing the scale of inhomogeneities. Therefore, the resolutions we find in our adiabatic simulations that are necessary to resolve structure in the unheated IGM are also sufficient to model radiative feedback at all times.  The outline of the paper is as follows. Details of the simulation setup and radiative transfer are described in $\\S 2$. In $\\S 3$ we present our numerical results, followed by $\\S 4$, where we present the results of our convergence tests. $\\S 5$ concludes with a discussion of our main results.  \\begin{table}[t] \\caption{Simulation Parameters} \\centering \\begin{tabular}{c c c c c c c c} \\hline\\hline Simulation & $N$ & $L$ (Mpc) & $m_{\\rm dm}$ ($M_\\odot$) & $r_{\\rm soft}$ (pc) \\\\ \\hline A1 & $2\\times 256^3$ & 0.25 & 31                 & 30  \\\\ A3 & $2\\times 256^3$ & 1    & $2.0 \\times 10^3$  & 120 \\\\ A4 & $2\\times 256^3$ & 2    & $1.6 \\times 10^4$  & 240 \\\\ A6 & $2\\times 256^3$ & 8    & $1.0 \\times 10^6$  & 960 \\\\ \\hline B1 & $2\\times 512^3$ & 0.25 & 3.8                & 15  \\\\ B2 & $2\\times 512^3$ & 0.5  & 31                 & 30  \\\\ B3 & $2\\times 512^3$ & 1    & 240                & 60  \\\\ B4 & $2\\times 512^3$ & 2    & $2.0 \\times 10^3$  & 120 \\\\ B5 & $2\\times 512^3$ & 4    & $1.6 \\times 10^4$  & 240 \\\\ B6 & $2\\times 512^3$ & 8    & $1.3 \\times 10^5$  & 480 \\\\ \\hline C1 & $2\\times 1024^3$ & 0.25 & 0.5               & 7.5 \\\\ C2 & $2\\times 1024^3$ & 0.5  & 3.8             & 15 \\\\ C3 & $2\\times 1024^3$ & 1    & 31                & 30  \\\\ C4 & $2\\times 1024^3$ & 2    & 240               & 60  \\\\ C6 & $2\\times 1024^3$ & 8    & $1.6 \\times 10^4$ & 240 \\\\ \\label{table:simparams} \\end{tabular} \\end{table} ",
        "conclusions": "We have performed high-resolution, cosmological simulations of structure formation at redshifts $z>6$, including adiabatic hydrodynamics. By post-processing the resulting density fields with a radiative transfer algorithm for hydrogen ionizing radiation, we have determined the opacity of the unheated IGM, in terms of the mean free path to ionizing radiation, $\\lambda$, as a function of redshift and ionizing background intensity. These results are relevant (1) as converged solutions for the opacity of the IGM early in the reionization process, before photoheating has evaporated small-scale structure and (2) in determining what mass and length resolutions are necessary to correctly model the propagation of ionization fronts into the neutral IGM. We derive values of $n_{\\rm crit}$, the proper hydrogen number density above which gas remains neutral, that are for the most part a function of only $\\Gamma_{-12}$. Simulations that mimic the effect of self-shielding by turning off the optically-thin flux at high densities should use $n_{\\rm crit} \\sim 0.1$~cm$^{-3}~\\Gamma_{-12}^{2/3}$, independent of redshift.  Our post-processing approach neglects the hydrodynamic feedback of photoheating on the density evolution. These results therefore indicate what the initial degree of inhomogeneity should be as ionization fronts propagate into the IGM. In addition, they place an upper limit to this inhomogeneity in patches of the IGM that have already been ionized. We find that the initial clumping factor of the IGM just as it is being ionized is a strong function of redshift and ionizing background intensity, with typical values at $z=10$ ranging from about $c_l=4.4$ to $16$ and $\\lambda=0.7$ to $15$~Mpc, for $\\Gamma_{-12}=0.03$ to $\\Gamma_{-12}=3$, respectively.  Modelling the transition from a neutral to ionized IGM requires self-consistent simulations of the coupled radiative transfer and hydrodynamical photoevaporation process. \\citet{shapiro/etal:2004} used idealized two-dimensional radiative transfer hydrodynamics calculations of the photoevaporation of initially spherical, isolated minihalos, surrounded by infalling gas. Those calculations showed that smaller minihalos are photevaporated faster, and that larger fluxes lead to faster photoevaporation times as well. \\citet{2005MNRAS.361..405I} extended these models to show that the typical timescale for minihalo photoevaporation is $t_{\\rm ev} \\sim 10-100$ Myr. This is comparable to the recombination time $t_{\\rm rec} = 1/(c_l \\alpha_B \\langle n_H \\rangle) \\sim 100$ Myr at $z \\sim 10$ for a clumpy IGM with $c_l = 10$. This suggests that Jeans smoothing of the IGM occurs before recombinations have had time to significantly disturb the reionization process. The amount by which recombinations occurring within minihalos could delay reionization was studied by \\citet{ciardi/etal:2006} who performed numerical simulations using the results of \\citet{2005MNRAS.361..405I} as a subgrid model for minihalo absorption. For the extreme case where Jeans smoothing fails to suppress minihalo formation they place $\\Delta z \\sim 2$ as the upper limit to the redshift delay of reionization induced by minihalos. For the opposite and more physically realistic case where minihalo formation is heavily suppressed by photoionization they find only a modest impact on reionization with a volume averaged ionized fraction that is $\\lesssim 15\\%$ lower than the case where minihalo recombinations are ignored.  However, the photoevaporative process is in reality likely to be more complex than for the simplified geometries and source lifetimes considered by \\citet{shapiro/etal:2004}, with halos over a range of masses clustered in space and arranged within a ``cosmic web'' of filamentary structure. Filamentary infall from nearby neutral gas could replenish halos as they are being evaporated, considerably extending the photoevaporation process, while ionization from highly luminous but intermittent starburst galaxies could result in large clumping factors and stalled minihalo evaporation, considerably increasing photon consumption and leading to a much more complex morphology of early H II regions \\citep[e.g.,][]{wise/etal:2012} than is typically envisioned.  A possible scenario that we have not considered here is the suppression of gas clumping at early times due to the presence of high-redshift X-ray sources \\citep[see, e.g.,][]{haiman:2011}. These may be associated with traditional X-ray sources like supernova or by more exotic sources like microquasars \\citep{mirabel/etal:2011}. X-rays have a small absorption cross-section meaning that a high-redshift distribution of X-ray sources would expose the IGM to a nearly uniform source of heating, inhibiting minihalo formation and growth at early times \\citep{oh/haiman:2003}. The result is a warm ($T \\sim 1000$ K) and weakly ionized IGM that would later become reionized by the patchy network of star-forming galaxies with softer radiation spectra. In this case the clumping factor would already be reduced at the onset of reionization and the resolution requirements presented here would become less strict. Our convergence criteria may therefore be considered the conservative case where the IGM has not been smoothed by heating processes prior to reionization.  We find that convergence is reached at a dark matter particle mass of $m_{\\rm dm} \\lesssim 50~ M_\\odot$. A box size of $L \\gtrsim 1$ Mpc is necessary to sample the IGM for the purpose of modelling absorptions by small-scale structure. The clumping factors we find from our converged results are somewhat smaller than the values $c_l\\sim 30$ found in early attempts to characterize the clumpiness of the IGM which did not accurately separate ionized and neutral gas \\citep[e.g.,][]{gnedin/ostriker:1997}, but are higher than the clumping factors found by \\citet{pawlik/etal:2009} at $z\\sim 9$, just before the IGM in those simulations was heated by ionizing radiation. We attribute this difference to the increased mass resolution of our simulations, which resolve halo masses down to the Jeans mass in an {\\em unheated} IGM ($\\sim 10^4 M_\\odot$), as opposed to that corresponding to the Jeans mass for a photoionized gas temperature of $\\sim 10^4$~K ($\\sim 10^8 M_\\odot$).  As pointed out by \\citet{pawlik/etal:2009}, the clumping factors they find at $z\\sim 6$, for a patch of the IGM which was ionized significantly earlier, at $z\\sim 9$, are converged with respect to the Jeans scale of the heated IGM. Their value likely approaches the {\\em correct}\\footnote{The correctness of the clumping factor depends on the specific physicalal processes affecting the evolution of baryons within the IGM and on the particular context to which the clumping factor is being used to describe. Here we use the term correct to refer to the value of the clumping factor that would be obtained if the simulation in question had infinite resolution.} value for the post-reionization IGM at $z = 6$ because a long enough time had passed since the gas was ionized for photoheating to evaporate existing small-scale structure and suppress accretion onto newly formed dark matter minihalos with masses below $\\sim 10^{8-9}~ M_\\odot$, which were resolved in their highest resolution simulations by $\\sim 100-1000$ dark matter particles.  \\begin{figure} \\smjustify \\includegraphics[width=\\smwidth]{fig10.pdf} \\vspace{-0.7cm} \\caption{Here we compare clumping factors $C_{-1}$ and $C_{100}$ from our fiducial simulation B2 ($512^3$, 0.5 Mpc) to the \\citet{pawlik/etal:2009} simulation L6N256 ($256^3$, 6 Mpc). The Pawlik et al. curve corresponds to their reference simulation that did not include photoionization and has similar attributes to our simulation A6 ($256^3$, 8 Mpc). As expected, A6 and L6N256 are in good agreement with each other while the higher-resolved simulation B2 shows clumping factors a few times larger than L6N256.} \\label{figure:pawlik} \\vspace{0.2cm} \\end{figure}  \\citet{pawlik/etal:2009} demonstrate convergence of their clumping factor for the heated IGM while noting that convergence would be more difficult to obtain for an unheated IGM. We have explicitly demonstrated this latter point here and accordingly show that their values for the clumping factor are likely underestimates during the initial stages of reionization, by about a factor of a few. This is illustrated in Figure \\ref{figure:pawlik} where we compare clumping factors from two of our simulations to the unheated simulation L6N256 of Pawlik et al. Here we plot the clumping factors $C_{-1}$ and $C_{100}$ with this notation being used to emphasize that these clumping factors are evaluated using all gas below some density cutoff.  For $C_{-1}$ a physical density threshold of $n_{\\rm crit} = 0.1~{\\rm cm}^{-3}$ is used while an overdensity threshold of $\\Delta_{\\rm crit} = 100$ is used for $C_{100}$. These definitions differ from the definition of $c_l$ in equation (\\ref{eq:cl}) that we have been using so far which involves an average over all ionized gas in the box. In any event, the clumping factors we find in our fiducial simulation B2 are at all times larger, by a factor of 1.2 at $z = 20$ and 3.5 at $z = 6$ for $C_{-1}$. As described above, the clumping factors we obtain at low redshift are overestimates, and in reality should be closer to $C_{-1} \\sim 6$ or $C_{100} \\sim 3$ found by Pawlik et al. in the presence of a photoevaporative background. Combining these two results, the clumping factor of the the IGM evolves strongly just after a patch of IGM is ionized. For ionization at $z=9$, the clumping factor drops from $c_l \\sim 20$ at $z=9$ to a few at $z=6$, depending on the intensity of the ionizing background -- with larger intensities leading to higher clumping factors and larger mean free paths. This strong suppression of the clumping factor due to photoheating was demonstrated by Pawlik et al. who referred to it as a positive feedback on reionization since it reduces the total number of recombinations occurring within small-scale absorption systems.  The results presented here for the inhomogeneity of electron density in the presence of an ionizing background should serve as a foundation for more detailed study of radiative transfer and hydrodynamical effects in the initial stages of reionization, including the effects of the initial relative velocity between baryons and dark matter \\citep[e.g.,][]{tseliakhovich/hirata:2010}, preheating by long mean free path X-ray photons \\citep[e.g.,][]{ricotti/ostriker:2004,ricotti/etal:2005}, and photoevaporation \\citep[e.g.,][]{shapiro/etal:2004, abel/etal:2007}. In simulating all of these processes, it will be necessary to resolve small-scale structure in the way outlined here."
    },
    "1208/1208.1356_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "\\end{table}"
    },
    "1208/1208.3509_arXiv.txt": {
        "abstract": " ",
        "introduction": "A scalar field is the simplest fundamental matter field that we may introduce in a cosmological model, since it adds essentially just one degree of freedom and it is naturally invariant under coordinate transformation. Even though, we can quote two different couplings of a scalar field: the minimal coupling, leading to the so-called Einstein's frame; the non-minimal coupling, corresponding to the Jordan's frame. Of course, the minimal coupling is the most simple one. But, the non-minimal coupling is one of the predictions of string theory: the scalar field emerging from this theory, called dilaton field, appears non-minimally coupled to the gravity term \\cite{copeland}. Moreover, being an interesting alternative to the General Relativity theory, the Brans-Dicke theory \\cite{bd}, is based on the non-minimal coupling. \\par Both frames, the Einstein and Jordan ones, may be connected by a conformal transformation. There is an intensive discussion in the literature about the meaning of such connection, in the sense that it may be just a mathematical mapping of one system into another, or may hide a deep physical meaning, see for example \\cite{frame} and references therein. For some observables, it is clear that the predictions in one frame are not equivalent to the predictions in the other frame. In quantum cosmology this problem has been treated in references \\cite{nelson1,nelson2}, showing that it is possible - even at quantum level - to mapp the equations in one frame into another frame, but with different predictions to the evolution of the universe. \\par In references \\cite{nelson1,nelson2} a system consisted of gravity and a scalar field, minimally or non-minimally coupled has been considered. In this case, the analysis present a major challenge, typical of quantum cosmology: how to obtain predictions for the evolution of the universe as function of time? In fact, it is well known that, considering the Einstein-Hilbert action, in quantizing it the resulting equation, the Wheeler-de Witt equation, has no time parameter due to the invariance of the gravitational system under time reparametrization. The introduction of a scalar field does not change this situation. In those references, the $WKB$ method has been used in order to obtain predictions for the evolution of the universe from the wavefunction determined from the (timeless) Wheeler-de Witt in the minisuperspace. Concerning this question, see reference \\cite{nelsonr} and references therein. \\par Another possibility to obtain the time evolution of a quantum cosmological system in the minisuperspace is to introduce matter in the form of a fluid, employing the Schutz description based in terms of potential that conveys the degrees of freedom of the fluid \\cite{schutz1,schutz2}. It has been shown that such description leads, always in the minisuperspace, to a Schr\\\"odinger-like equation, with the degrees of freedom of the matter playing the r\\^ole of time. This proposal has been presented in reference \\cite{rubakov}, and extensively analysed in reference \\cite{lemos}. \\par Our aim in this work is to address the problem of time evolution of the universe in the presence of a dilaton-like field from a quantum cosmological perspective. Since we are interested mainly in the primordial universe, a radiative fluid will be introduced. This may allow to recover the time variable through the employment of the Schutz formulation. Performing a conformal transformation (which does not affect the radiative fluid, since it is conformal invariant), we re-write the dilaton-gravity system in the Einstein's frame, from which the Wheeler-de Witt equation in the minisuperspace is constructed, resulting in a Schr\\\"odinger-like equation. From this Schr\\\"odinger-like equation, an explicit solution is obtained by using a specific ordering factor. A wavepacket is determined, but its norm is time-dependent. Hence, it does not fit in the usual many-worlds interpretation of quantum mechanics \\cite{tipler,nelsonr}, but it admits a Bohmian analysis \\cite{nelsonr,holland}. The Bohmian trajectories contain universes that are singurality-free. \\par This paper is organized as follows. In the next section, we quantize the dilaton-gravity-radiation system in the mini-superspace. In section 3, a wavepacket is constructed and a (formal) many-world analysis is performed. In section 4 the Bohmian trajectories are studied. In section 5 we present our conclusions. ",
        "conclusions": "The possibility to make predictions and subsequently extract a family of solutions (trajectories in the minisuperspace) is a recurrent problem in quantum cosmology. Although the well-known $WKB$ method, including decoherence features, has provided a vast range of particular results, there is an alternative framework by means of transforming the $WdW$ equation into a Schr\\\"odinger equation, with the time variable induced by a matter component \\cite{schutz1,schutz2,rubakov}. Quite interesting results have appeared in the literature \\cite{lemos} using this procedure. \\par In this paper, we investigated how we could obtain some predictions with that alternative framework, using a system characterized by a dilaton field and radiation expressed using the Schutz variables. The dynamics of the radiative fluid implies a time variable. \\par We found that, in order the Schr\\\"odinger equation to be elliptic, leading to a positive energy spectrum, the dilaton field must have a phantom behaviour. On the other hand, the construction of a quasi gaussian superposition, conjugated with some convergence criteria, lead to a wavepacket with a time dependent norm and hence unitarity could not be invoked. Nevertheless, this still allows us to investigate this wavepacket under a Bohmian perspective. Interestingly, we found a bouncing behaviour for the scale factor, i.e., the singularity is avoided. In this sense we extend the results of reference \\cite{brasil}, where the imposition of unitarity led to a conclusion that the dilaton field should be constant. At same time, such approach open new perspective concerning anisotropic quantum models \\cite{brasilbis}. \\par There are still many open directions to explore in this program. In particular, we must look for, (i) other wavepackets, presumably by using numerical methods, (ii) study the dynamics in the Jordan frame. We hope to address these problems in a future work.  {\\bf Acknowledgements:} We thank CNPq for partial financial support. PVM is grateful to CENTRA-IST for financial assistance. He also wants to thank UFES, where this work was completed, for hospitability"
    },
    "1208/1208.0600_arXiv.txt": {
        "abstract": "We study the cosmology of Galileon modified gravity models in the linear perturbation regime. We derive the fully covariant and gauge invariant perturbed field equations using two different methods, which give consistent results, and solve them using a modified version of the {\\tt CAMB} code. We find that, in addition to modifying the background expansion history and therefore shifting the positions of the acoustic peaks in the cosmic microwave background (CMB) power spectrum, the Galileon field can cluster strongly from early times, and causes the Weyl gravitational potential to grow, rather than decay, at late times. This leaves clear signatures in the low-$l$ CMB power spectrum through the modified integrated Sachs-Wolfe effect, strongly enhances the linear growth of matter density perturbations and makes distinctive predictions for other cosmological signals such as weak lensing and the power spectrum of density fluctuations. The quasi-static approximation is shown to work quite well from small to the near-horizon scales. We demonstrate that Galileon models display a rich phenomenology due to the large parameter space and the sensitive dependence of the model predictions on the Galileon parameters. Our results show that some Galileon models are already ruled out by present data and that future higher significance galaxy clustering, ISW and lensing measurements will place strong constraints on Galileon gravity. ",
        "introduction": "The accumulated evidence for the present-day accelerated expansion of the Universe, driven by what is generically referred to as `dark energy', is now overwhelming \\cite{Reid:2009xm, Komatsu:2010fb, Amanullah:2010vv}. The simplest explanation for the nature of dark energy is a simple cosmological constant but, despite the good agreement with the observational data so far, such an explanation is plagued with serious fine tuning and coincidence problems. This has motivated the proposal of alternative models to explain the observations, the majority of which fall into two classes. The first one assumes the existence of a dynamical dark energy field (often of scalar type) which dominates the energy density today and has a negative pressure to accelerate the Universe \\cite{2006IJMPD..15.1753C, Li:2011sd}. The other considers that the standard law of gravity, general relativity, fails on cosmological scales and must be completed by modifications capable of accelerating the Universe \\cite{Clifton:2011jh}. Models in the second class have attracted a lot of research interest recently, and significant progress has been made in both the theoretical modelling \\cite{Brax:2011aw, Brax:2012gr} and numerical simulations \\cite{Li:2011vk,Jennings:2012pt,Li:2012by}.  One notable example of a modified gravity model which has been the subject of many recent papers is the Galileon model \\cite{PhysRevD.79.064036,PhysRevD.79.084003}. Here, the deviation from general relativity is mediated by a scalar field $\\varphi$, dubbed the Galileon, whose Lagrangian is invariant under the Galilean shift symmetry $\\partial_\\mu\\varphi \\rightarrow \\partial_\\mu\\varphi + b_\\mu$ (hence the name), where $b_\\mu$ is a constant vector. Such a field appears, for instance, as a brane-bending mode in the decoupling limit of the four-dimensional boundary effective action of the DGP braneworld model \\cite{2000PhLB..485..208D, 2003JHEP...09..029L, 2004JHEP...06..059N} which was proposed well before the Galileon model. However, in spite of being theoretically appealing, the self-accelerating branch of the DGP model, which is of interest to the cosmological community, is plagued by the ghost problem \\cite{1126-6708-2003-09-029, 1126-6708-2004-06-059, PhysRevD.73.044016} (i.e. there is not a well defined minimum energy). Taking the DGP model as inspiration, it was shown in \\cite{PhysRevD.79.064036} that in four-dimensional Minkowski space there are only five Galilean invariant Lagrangians that lead to second-order field equations, despite containing highly nonlinear derivative self-couplings of the scalar field. The second-order nature of the equations of motion is crucial to avoid the presence of Ostrogradski ghosts \\cite{Woodard:2006nt}. In \\cite{PhysRevD.79.084003, Deffayet:2009mn}, it was shown how these Lagrangians could be generalised to curved spacetimes. These authors concluded that explicit couplings between the Galileon field derivatives and curvature tensors are needed to keep the equations of motion up to second-order. Such couplings however break the Galileon symmetry which is only a symmetry of the model in the limit of flat spacetime \\cite{Hui:2012qt}. The couplings of the Galileon field and the curvature tensors in the equations of motion change the way in which spacetime responds to matter distributions, which is why the Galileon model is a subclass of modified gravity theories.  Since the equations of motion are kept up to second order, it means that the Galileon model is a subclass of the more general Horndeski theory \\cite{Horndeski:1974wa, Deffayet:2011gz, Kobayashi:2011nu}. The Horndeski action is the most general single scalar field action one can write that yields only second order field equations of motion of the metric and scalar fields. Besides the Galileon model, it therefore encompasses simpler cases such as Quintessence, k-essence \\cite{2006IJMPD..15.1753C} and $f(R)$ \\cite{Sotiriou:2008rp, DeFelice:2010aj} models as well as other models which also involve derivative couplings of the scalar field that have recently generated some interest such as Kinetic Gravity Braiding \\cite{Deffayet:2010qz, Pujolas:2011he, Kimura:2011td}, Fab-Four \\cite{Charmousis:2011bf, Charmousis:2011ea, Bruneton:2012zk, Copeland:2012qf, Appleby:2012rx}, k-mouflage \\cite{Babichev:2009ee} and others \\cite{Kobayashi:2009wr, Kobayashi:2010wa, Leon:2012mt}. An important difference between the Galileon model and some other corners of Horndeski's general theory is that in the Galileon model there are no free functions since the fuctional form of the Lagrangian is fixed by the shift symmetry (see however \\cite{DeFelice:2010nf}).  In any viable modified gravity theory, it is crucial that deviations from standard gravity get suppressed (or screened) in high matter-density regions where general relativity has been tested to high accuracy \\cite{1990PhRvL..64..123D, 1999PhRvL..83.3585B}. In the case of Galileon gravity, such a screening is realised via the Vainshtein mechanism \\cite{Vainshtein1972393}, which relies on the presence of the nonlinear derivative self-couplings of the Galileon field. Here, far away from gravitational sources, the nonlinear terms are subdominant and the Galileon field satisfies a linear Poisson equation (as the Newtonian potential), so that the extra (fifth) force mediated by it can be sizeable and proportional to standard gravity, effectively renormalising Newton's constant. Near the sources, on the other hand, the nonlinear terms become important, which strongly suppress the spatial variations of the Galileon field compared to that of the Newtonian potential and ensure that the extra force, which is the gradient of the Galileon field, is not felt on scales smaller than a given `Vainshtein radius'.  In certain respects, this is very similar to the chameleon screening \\cite{PhysRevD.69.044026, PhysRevD.70.123518}, which operates for instance in $f(R)$ gravity models \\cite{Sotiriou:2008rp, DeFelice:2010aj, Carroll:2004de, Brax:2008hh}. However, in the chameleon case the self-interaction of the scalar field depends on the field value (through a nonlinear interaction potential) rather than its derivatives, and the non-derivative coupling of the scalar field to matter makes its behaviour highly sensitive to the environmental matter density -- in high density regions the field value, rather than merely its gradient, becomes extremely small so that the extra force is suppressed.  It is therefore evident that one has to go beyond the local environment to look for possible deviations from general relativity and distinct signatures of the different modified gravity models. In particular, a promising way is to look at the cosmic expansion and the formation of structure in the Universe: different screening mechanisms in different modified gravity models can lead to very different predictions as to when, where and how the various cosmological observables are affected.  The effects of Galileon gravity models on the background cosmological expansion have already been studied in the literature in great detail \\cite{Gannouji:2010au, PhysRevD.80.024037, DeFelice:2010pv, Nesseris:2010pc, Appleby:2011aa, PhysRevD.82.103015}. It has been shown that in these models there is  a stable de Sitter point that can be reached after the radiation and matter dominated eras, thus yielding a viable cosmological expansion history. Conditions to avoid the ghosts and other theoretical instabilities have also been derived by considering the linear perturbations \\cite{DeFelice:2010pv, Appleby:2011aa}.  To improve our understanding of the cosmological effects of Galileon gravity models and make direct comparisons with observational data, a proper investigation of the evolution of density fluctuations and formation of large-scale structure is necessary. Here, as an initial step, we consider the regime in which the density fluctuations are small such that their evolution is well described by linear perturbation theory. This regime is relevant for several important cosmological observables, such as the power spectrum of the cosmic microwave background (CMB) temperature fluctuations and its polarisations, the growth of matter density perturbations, the weak gravitational lensing of distant galaxies and the CMB map, and the integrated Sachs-Wolfe (ISW) effect and its cross correlation with the galaxy distribution. The rich information contained in this regime therefore warrants a detailed study of the Galileon effects, which is precisely the topic of this paper. The nonlinear regime of structure formation can in principle contain further interesting information, but its study is beyond the scope of the current paper.  The layout of this paper is as follows. We start by briefly presenting the Galileon model and the Galileon and metric field equations of motion in Section~\\ref{The model}. The perturbation equations are derived and presented in a covariant and gauge invariant (CGI) way in Section \\ref{The perturbation equations} using the method of $3+1$ decomposition. In Appendix~\\ref{alternative_derivation} we present an alternative and considerably simpler derivation of the perturbation equations which is particularly suitable for the Galileon model as it takes advantage of the fact that the Lagrangian density is fixed by the Galilean shift invariance and that there are no derivatives higher than second order. We present and discuss the results for the CMB, lensing and linear matter power spectra in Section \\ref{Results} which we obtain using a version of the {\\tt CAMB} code \\cite{camb_notes} which we have modified. In Section \\ref{Results} we also discuss the time evolution of the gravitational potential, Galileon field perturbation and Galileon density contrast and the validity of the quasi-static limit. We conclude in Section \\ref{Conclusion}.  Throughout this paper we will use the unit $c=1$ and metric convention $(+,-,-,-)$. Greek indices run over $0,1,2,3$ and we will use $8\\pi G=\\kappa=M^{-2}_{\\rm Pl}$ interchangeably, where $G$ is Newton's constant and $M_{\\rm Pl}$ is the reduced Planck mass. ",
        "conclusions": "We have studied the cosmology of Galileon gravity models at the linear perturbation level. For this we derived the full CGI perturbation equations using two independent methods: the normal procedure of linearising the full field equations and an alternative derivation that is particularly suitable for models like Galileon gravity, where the shape of the Lagrangian is fixed by certain symmetries (e.g., there are no free functions such as the potential in quintessence and $f(R)$ gravity models) and the field equations only contain up to second-order derivatives. The second derivation is particularly appealing because it is much simpler than the first one, which is very lengthy and complicated for the full Galileon model. We checked that the two methods give the same set of perturbation equations, and then solved these equations using a modified version of the {\\tt CAMB} code, which we tested by performing several successful consistency tests.   Our code also solves the background expansion history in Galileon models and our results agree with those in the literature. We find that the expansion rate in Galileon cosmology can depend sensitively on the initial value of the Galilean energy density, especially if the latter is not small, e.g., if $\\rho_{\\varphi, i} / \\rho_{m,i} \\gtrsim 10^{-5}$. Throughout the evolution, the expansion rate can be faster or slower than in $\\Lambda$CDM and the Galileon equation-of-state parameter can cross the phantom line ($w<-1$) in a way which is free of ghost-like instabilities.  The modified background expansion translates into a different age of the Universe and distance to last scattering, which leads to a visible shift in the positions of the acoustic peaks of the CMB temperature power spectrum. The strongest effect of the Galileon field on the CMB temperature power spectrum, however, appears to be on the largest angular scales (low values of $l$), where the full power receives a significant contribution from the integrated Sachs-Wolfe effect, which is due to the late-time evolution of the gravitational potential $\\phi$. Indeed, we found that in Galileon models the gravitational potential evolves even during the matter dominated era and can undergo an overall deepening at late times. This is very different from the standard $\\Lambda$CDM prediction that the gravitational potential is constant during matter domination and becomes shallower when the expansion of the Universe starts to accelerate. The origin of the abnormal evolution in the gravitational potential can be traced back to the pressure perturbation and anisotropic stress of the Galileon field, which cause it to cluster strongly (comparable to the clustering of dark and baryonic matter species) on all scales.  The evolution of the gravitational potential influences a number of cosmological observables, both directly and indirectly. In addition to the ISW effect, it also has strong impact on the growth of matter density perturbations (and therefore the linear and nonlinear matter power spectra), weak gravitational lensing and their cross correlations. In particular, we have shown that the Galileon model can predict considerably more power than $\\Lambda$CDM for the weak lensing power spectrum at all scales, even if their predictions for the CMB power spectrum more or less agree. Galileon models might also have distinctive predictions for the cross correlation of the ISW effect with the galaxy distribution. These are important observational signatures in the linear perturbation regime that can in principle help to distinguish the Galileon models from the standard $\\Lambda$CDM paradigm.   On the other hand, the sensitive dependence of the Galileon behaviour on the model parameters makes the phenomenology of the Galileon cosmology especially rich. For example, by tuning the parameters in the Galileon Lagrangian, one can get a CMB power spectrum which is very close to the $\\Lambda$CDM prediction and therefore hard to distinguish by looking at very large scales.  On subhorizon scales, we have seen that the linear growth of matter density perturbations  can be significantly enhanced with respect to the $\\Lambda$CDM results, even for those model parameters that lead to similar CMB power spectrum. However, in Galileon models, the Vainshtein screening mechanism is at play and its potential influence on the clustering of matter is still to be properly understood. As an analogy, in other modified gravity models such as the $f(R)$ gravity, the chameleon screening effect has been shown to make the linear perturbation theory a poor approximation even on scales as large as $k = 0.01 \\ \\rm{hMpc}^{-1}$. We therefore conclude that a better understanding of the true impact of the Vainshtein screening is necessary, before attempting a more rigorous confrontation of the predicted matter power spectrum with measurements of galaxy clustering. Such a study will be left for future work.  Finally, we have seen that the quasi-static approximation for the evolution of the Galileon field perturbation serves as a good approximation on subhorizon scales for the models we have shown in this paper. It works reasonably well on near-horizon scales such as $k = 0.001 \\ \\rm{hMpc}^{-1}$, with an error of the order of a few percent. However, for accuracy considerations we solve the full evolution equation of the Galileon perturbation in our code, which does not take much longer anyway.  In conclusion, we have shown that the detailed study of the full perturbation equations unveils a rich phenomenology in Galileon gravity models. The full cosmological parameter space increases considerably in Galileon gravity. Even with current data, the indications are that strong constraints can be placed on this parameter space. In a future project we will use our modification to the {\\tt CAMB} software to carry out a formal study of the Galileon parameter space."
    },
    "1208/1208.0560.txt": {
        "abstract": "{%Context The atmosphere of hot Jupiters can be probed by primary transit and secondary eclipse spectroscopy. Due to the intense UV irradiation, mixing and circulation, their chemical composition is maintained out of equilibrium and must be modeled with kinetic models.} {%Aims Our purpose is to release a chemical network, and the associated rate coefficients, developed for the temperature and pressure range relevant to hot Jupiters atmospheres. Using this network, we study the vertical atmospheric composition of the two hot Jupiters (HD~209458b and HD~189733b) with a model that includes photolyses and vertical mixing and we produce synthetic spectra.} {%Methods The chemical scheme is derived from applied combustion models that have been methodically validated over a range of temperatures and pressures typical of the atmospheric layers influencing the observations of hot Jupiters. We compare the predictions obtained from this scheme with equilibrium calculations, with different schemes available in the literature that contain N-bearing species and with previously published photochemical models.} {%Results Compared to other chemical schemes that were not subjected to the same systematic validation, we find significant differences whenever non-equilibrium processes take place (photodissociations or vertical mixing). The deviations from the equilibrium, and thus the sensitivity to the network, are more important for HD~189733b, as we assume a cooler atmosphere than for HD~209458b. We found that the abundances of \\ce{NH3} and HCN can vary by two orders of magnitude depending on the network, demonstrating the importance of comprehensive experimental validation. A spectral feature of \\ce{NH3} at 10.5~$\\mu$m is sensitive to these abundance variations and thus to the chemical scheme.} {%Conclusions Due to the influence of the kinetics, we recommend the use of a validated scheme to model the chemistry of exoplanet atmospheres. The network we release is robust for temperatures within 300-2500~K and pressures from 10 mbar up to a few hundreds of bars, for species made of C, H, O, N. It is validated for species up to 2 carbon atoms and for the main nitrogen species (\\ce{NH3}, HCN, \\ce{N2}, NO$_X$). Although the influence of the kinetic scheme on the hot Jupiters spectra remains within the current observational error bars (with the exception of \\ce{NH3}), it will become more important for atmospheres that are cooler or subjected to higher UV fluxes, departing more from equilibrium.} ",
        "introduction": "So far, more than 700 exoplanets have been confirmed and thousands of transiting candidates have been identified by the space telescope Kepler \\citep{2012arXiv1202.5852B}. Among them, hot Jupiters are a class of gas giants with orbital periods of a few days or less. They are found around $\\sim 0.5\\%$ of KGF stars \\citep{howard2010occurrence, howard2011}. About $10\\%$ of them transit their host star and their atmospheric composition and physical structure can be studied by transit spectroscopy \\citep[e.g.][]{charbonneau2000detection, charbonneau2008broadband, richardson2007spectrum, tinetti2007water, sing2008hubble, swain2008mid, swain2008presence, swain2009water, swain2009molecular, 2012arXiv1202.4721H}.  Although current observations are still limited and subjected to divergent interpretations, future instruments such as E-ELT, JWST \\citep{Gardner2006}, EChO \\citep{tinetti-Echo}, FINESSE \\citep{swain2010finesse} should provide better constraints on both the chemical composition and the temperature profiles of the nearby hot Jupiters like HD~189733b and HD~209458b. They will also be able to study more distant targets and deliver statistically significant trends about the nature of their atmospheres. Chemical modeling will be an important component of these studies. It will point to key observations able to distinguish between various hypotheses and will be used to analyze the observations and constrain, for instance, the atmospheric elemental abundances.  The first models of hot Jupiter atmospheres assumed chemical equilibrium \\citep[e.g.][]{burrows1999chemical, seager2000theoretical, sharp2007atomic, 2007ApJ...661L.191B, burrows2007theoretical, burrows2008theoretical, fortney2008unified}. However, strongly irradiated atmospheres are unlikely to be at chemical equilibrium. Their intense UV irradiation (typically 10,000 times the flux received on the top of the atmosphere of Jupiter) and strong dynamics result in photolyses and diffusion/advection timescales that are comparable or shorter than the chemical ones.  Deviations from the thermodynamic equilibrium have been discussed with timescale arguments \\citep[e.g.][]{lodders2002atmospheric, fortney2006influence, fortney2008synthetic, visscher2006, visscher2010, madhusudhan2010inversion}, or modeled with a few reactions describing the CO-\\ce{CH4} conversion coupled with the dynamics \\citep{cooper2006dynamics}. A more detailed modeling requires the use of a photochemical kinetic network. A kinetic network is, in practice, a list of reactions and associated rate coefficients able to describe quantitatively (within a certain accuracy) the kinetics of a pool of species, usually the most abundant ones. Constructing such network of reactions implies to solve two major questions. One has to do with the completeness of the network: What are the species and the reactions connecting them that must be included? The other issue is the availability of the kinetic data: The literature and databases may not provide the rate coefficients for some of the needed reactions or may provide conflicting values with no recommendation. These two issues are tightly connected and both depend on the considered range of temperatures and pressures. Eventually, and whatever the methodology adopted to select the reactions and their rates, it is the ability to predict experimentaly-controled abundances that can validate or not the network.\\\\ To investigate the consequences of the strong UV incident flux on neutral species, photochemical models have then been developed \\citep{liang2003source, liang2004insignificance}. Based on kinetics model dedicated to Jupiter's low-temperature atmosphere, these models however neglected endothermic reactions, which are in fact rather efficient in such hot atmospheres. \\citet{line2010high} introduced some endothermic reactions to a similar Jovian photochemical scheme but most of the pre-existing reactions were not reversed. They were therefore not able to reproduce the thermodynamic equilibrium, which occurs in the deep atmospheres of hot Jupiters. \\citet{zahnle2009atmospheric, zahnle2009soots}  developed a photochemical model considering the reversal of their whole set of two-bodies exothermic reactions. They selected their rate constants in the NIST database\\footnote{http://www.kinetics.nist.gov/kinetics} based on the following criteria: relevance of temperature conditions, date of review, date of the experiment and date of the theoretical study (in order of preference). \\citet{moses2011disequilibrium} developed a model that considered the reversal of all the reactions, including three-body reactions, ensuring to reproduce the thermodynamical equilibrium from the top to the deepest layers of the atmosphere. Their chemical scheme is derived from the Jupiter and Saturn models \\citep{gladstone1996hydrocarbon, moses1996, moses1995post, moses1995nitrogen, moses2000photochemistry, moses2000photochemistry2} with further updates on the basis of combustion-chemistry literature. None of the aforementioned works discuss the validation of the chemical scheme against experiments. In addition, the fact that computed abundances evolve towards the composition predicted by equilibrium calculations (at given pressure \\textit{P} and temperature \\textit{T} with no external irradiation nor mixing) is by no means a validation of the kinetic network. Indeed, any network containing at least as many independent reversible reactions as modeled species, in which the rates for the backward processes are derived from equilibrium constants and forward rates, will evolve toward the equilibrium predicted with the same equilibrium constants, whatever the quantitative values of the forward rates, as illustrated in Fig.~\\ref{fig:0D_GRIM_NOM}. \\begin{figure}[!htbp] \\centering {\\includegraphics[width=0.95\\columnwidth]{images/Resultats_NH3_o_comp_nom_grim.eps}} \\caption{Abundances of \\ce{NH3} as a function of time computed with two kinetic schemes fully reversed according to equilibrium constants but differing by their nitrogen chemistry (nominal and \\textit{GRIMECH} as defined in Sects.~\\ref{sec:nitrogen} and \\ref{sec:nitrogen2}). While they both converge towards the equilibrium (dotted line) they exhibit very different evolution. Initial condition is a mixture of \\ce{H2}, \\ce{CH4}, \\ce{O2}, \\ce{N2} and \\ce{He} with solar elemental abundances.} \\label{fig:0D_GRIM_NOM} \\end{figure} Fortunately, and due to the physical conditions and elemental composition of hot Jupiters (and hot Neptunes) atmospheres, we benefit from decades of intensive work done in the field of combustion, that includes a vast amount of experiments, the development of comprehensive mechanisms\\footnote{In the field of combustion, a \\textit{mechanism} or {reaction base} is a network of reactions able to describe the kinetic evolution of a given pool of species. The mechanism includes the list of reactions and the associated rate coefficients, in a modified Arrhenius form, as well as the thermodynamic data for all the species involved in these reactions, which are required to calculate the equilibrium constants of the reactions and the rates of the reverse reactions} and the systematic comparison between the two. Therefore, we propose in the present work a new mechanism dedicated to the chemical modeling of hot atmospheres that is not adapted from previous Solar System photochemical models but derives from industrial applications (mainly combustion in car engines). Details about this chemical network and its range of validity are presented in Sect.~\\ref{sec:kinetic}.  We use this chemical network in a 1D model that includes photolyses and vertical transport, which has been previously used to study the atmospheric photochemistry of various objects of the Solar System: Neptune \\citep{dobrijevic2010key}, Titan \\citep{hebrard2006photochemical,hebrard2007photochemical}, Saturn \\citep{dobrijevic2003effect, cavalie2009first}, and Jupiter \\citep{cavalie2008observation}) as well as extrasolar terrestrial planets \\citep{selsis2002signature}. We model the photochemistry of two hot Jupiters: HD~209458b and HD~189733b (Sect.~\\ref{Sec:application}). We study the departure from thermodynamic equilibrium and compare our results with those of \\citet{moses2011disequilibrium} (Sect.~\\ref{sec:results}). We also investigate how including different reaction networks specific to nitrogen-bearing species influences the model results (Sect.~\\ref{sec:nitrogen2}) and the planetary synthetic spectra (Sect.~\\ref{sec:spectra}). ",
        "conclusions": "photodissociations do not have a significant impact on the atmospheric composition of HD~209458b, with the high temperatures we assume. It remains at the thermodynamic equilibrium for pressures higher than 1 bar. For lower pressures vertical transport affects the abundances of HCN, \\ce{NH3} and \\ce{CH4} and some of the minor species associated. For HD~189733b,  we assume significantly lower temperatures and find the atmospheric composition to be more sensitive to photolyses and vertical transport, all species being affected, except the main reservoirs, \\ce{H2}, \\ce{H2O}, CO and \\ce{N2}. Quantitatively, however, we find significant differences (up to several orders of magnitude in the case of HD~189733b) in the abundances that are likely to be due to the different chemical schemes used. These differences are smaller for HD~209548b because kinetics have less influence. The quenching of HCN and \\ce{NH3}, as well as \\ce{CH4} to a lower extent, is particularly affected, as well as most species sensitive to photochemistry in the upper atmosphere. Despite being large in terms of abundances, these differences do not produce strong effects on the spectra, with the exception of \\ce{NH3} at 10.5~$\\mu$m. Confronting different schemes with observations will thus have to wait more accurate spectroscopic observations (JWST, E-ELT, EChO). Until that, experimental validation appears mandatory.\\\\ In order to illustrate the sensitivity to the kinetic scheme, we implemented different available nitrogen schemes, that are either optimized (\\textit{GRIMECH}, \\textit{GDFKin}) or non validated (\\textit{Dean}). We studied the extent of the possible results, and found large differences whenever disequilibrium chemistry is at work. Changing the nitrogen scheme strongly affects the quenched species (HCN, \\ce{NH3}) and most species (including hydrocarbons) in the upper atmosphere of HD~189733b. For HD~209458b, deviations are again less noticeable as the atmosphere departs less from equilibrium.  We therefore emphasize on the need to use validated and non optimized chemical schemes. This is already true for hot Jupiters but this is even more crucial in the case of cooler atmospheres (GJ1614b, GJ3470b, for instance), which depart more from thermodynamic equilibrium and are more sensitive to kinetics.\\\\ Our nominal scheme can be downloaded from the KIDA database\\footnote{http://kida.obs.u-bordeaux1.fr/} \\citep{wakelam_kida2012}. The scheme is designed to reproduce the kinetic evolution of species with less than two carbon atoms. In order to study atmospheres with C/O ratio higher than solar (close to or above 1), we are currently developing a C$_0$-C$_6$ scheme, which will be able to describe kinetics of species up to 6 carbon atoms. One of the next improvement of our model should be the addition of sulfur. As kinetics of nitrogen species is an active field of research, we expect regular updates of the network (which would be notified and available on KIDA).  Note also that conclusions of this study on the chemical composition of hot Jupiters, which derive from models using an average 1D vertical profile, will probably have to be revisited with the effects of atmospheric circulation."
    },
    "1208/1208.5480_arXiv.txt": {
        "abstract": "{ 176P/LINEAR is a member of the new cometary class known as main-belt comets (MBCs). It displayed cometary activity shortly during its 2005 perihelion passage, which may be driven by the sublimation of subsurface ices.  We have therefore searched for emission of the \\ce{H2O} \\trans{} ground state rotational line at 557 GHz toward 176P/LINEAR with the Heterodyne Instrument for the Far Infrared (HIFI) onboard the \\herschel{} Space Observatory on UT 8.78 August 2011, about 40 days after its most recent perihelion passage, when the object was at a heliocentric distance of 2.58 AU.  No \\ce{H2O} line emission was detected in our observations, from which we derive sensitive 3-$\\sigma$ upper limits for the water production rate and column density of $<~\\upper \\times 10^{25}$ \\s\\ and of $< \\column \\times 10^{10}$ cm$^{-2}$, respectively.  From the peak brightness measured during the object's active period in 2005, this upper limit is lower than predicted by the relation between production rates and visual magnitudes observed for a sample of comets at this heliocentric distance.  Thus, 176P/LINEAR was most likely less active at the time of our observation than during its previous perihelion passage.  The retrieved upper limit is lower than most values derived for the \\ce{H2O} production rate from the spectroscopic search for CN emission in MBCs.} ",
        "introduction": "Classical main-belt asteroids are small bodies that orbit the Sun in low inclination and low eccentricity orbits between the orbits of Mars and Jupiter.  Physically, asteroids are thought to be devoid of volatiles, while comets are icy bodies that become active in the inner solar system thanks to the sublimation of ices, mostly water.  Comets originate in the outskirts of the solar system beyond the snow line, where temperatures in the solar nebula were low enough for water to condense onto icy grains \\citep{1981PThPS..70...35H}.  A new class of bodies has been discovered recently, the so-called main-belt comets (MBCs), which have orbital properties that are indistinguishable from standard asteroids with a Tisserand parameter with respect to Jupiter that is greater than three, and they display cometary activity in the form of a dust tail during part of their orbit.  Numerical simulations have shown that these objects are not comets from the Kuiper Belt or Oort Cloud that have been recently transferred to orbits within the main belt, but instead are most likely formed in situ at their current locations \\citep{2002Icar..159..358F}.  Theoretical models suggest that the snow line was initially close to the Mars orbit due to the absorption of stellar radiation by dust \\citep{2000ApJ...528..995S,2006ApJ...640.1115L}. Therefore, objects formed at their current locations in the outer asteroid belt may have been able to accumulate water ice in subsurface reservoirs, despite the effect of solar radiation.  Determining the composition of this class of objects can provide important clues to both the thermal properties that allow water to survive in subsurface layers and the distribution of volatile materials in the solar nebula to constrain planet formation mechanisms. Additionally, MBCs may have played an important role in the delivery of water and other volatiles to the inner solar system, including the Earth.  The MBC 176P/LINEAR (hereafter 176P) was discovered in 1999 and originally categorized as asteroid 118401 LINEAR. This object belongs to the Themis asteroid family. Cometary activity was reported for this object around perihelion in 2005 \\citep{2011AJ....142...29H} by the Hawaii Trails project \\citep[HTP;][]{2006Sci...312..561H,2009A&A...505.1297H}.  It displayed a mean photometric excess of $\\sim30$\\% during a month-long active period around its perihelion passage, consistent with an approximate total dust mass-loss of $\\sim 7\\times10^4$\\ kg \\citep{2011AJ....142...29H}.  Although ice sublimation is expected to trigger MBC activity, gas emission has never been directly detected in these objects owing to their low activity, which requires very sensitive observations.  {\\it Herschel} proves to be the most sensitive instrument for directly observing water in a distant comet \\citep[e.g.][]{2010DPS....42.0304B}.  In this paper we present the \\herschel{} observation of the \\trans{} fundamental rotational transition of \\ce{H2O} at 557 GHz in 176P.  This observation is intended to test the prediction that the observed cometary activity of MBCs is driven by sublimation of water ices and to constrain the production process. ",
        "conclusions": "\\begin{table} \\caption{Standard deviation of the brightness temperature and line area, and retrieved 3-$\\sigma$ upper limits of the \\ce{H2O} production rate in 176P/LINEAR.} \\label{tbl:q} \\centering \\begin{tabular}{c c c c} \\hline\\hline Spectrometer & $\\sigma_{T_\\mathrm{mB}}$ & $\\sigma_{\\int T_\\mathrm{mB}\\, dv}$ & $Q_\\ce{H2O}$\\tablefootmark{a} \\\\ & (K) & (K \\kms) & (\\s) \\\\ \\hline WBS & $6.598 \\times 10^{-4}$ & $6.172 \\times 10^{-4}$ & $< 2.08 \\times 10^{25}$ \\\\ HRS & $1.998 \\times 10^{-3}$ & $6.365 \\times 10^{-4}$ & $< 2.14 \\times 10^{25}$ \\\\ \\hline \\end{tabular} \\tablefoot{ \\tablefoottext{a}{Production rates derived for a gas kinetic temperature of 20 K, expansion velocity of 0.5 \\kms, and an electron density scaling factor of $\\xne = 0.2$.} } \\end{table}  We observed water emission in 176P with \\herschel/HIFI to test the prediction that cometary activity in MBCs is driven by sublimation of water ices from the nucleus.  There are several mechanisms that have been proposed to drive mass loss from small bodies, including sublimation of subsurface ices, rotational instability, impact ejection and thermal fracture \\citep[see][for recent reviews of MBC physical properties and activation mechanisms]{2011P&SS...59..365B,2012AJ....143...66J}.  The cometary activity observed in 176P was initially found to suggest the presence of sublimating subsurface ice that may have been exposed by recent collisions \\citep{2011AJ....142...29H}.  This view is supported by the detection of water ice absorption in spectroscopic observations centered at 3.1 $\\mu$m of the surface of the largest asteroid of the Themis asteroid family, 24 Themis, which belongs to the same dynamical family as 176P \\citep{2010Natur.464.1322R,2010Natur.464.1320C}, although it has been claimed that the measured spectra are consistent with the transmission spectra of goethite \\citep{2011A&A...526A..85B}.  From the search for the \\ce{H2O} \\trans{} rotational line at 557 GHz in 176P, a 3-$\\sigma$ upper limit for the \\ce{H2O} production rate of $< \\upper \\times 10^{25}\\ \\s$ is derived from the WBS and HRS data, for gas expansion velocities between 0.4--0.7 \\kms and gas kinetic temperatures between 20 and 40 K.  Using the peak value of the $R$-band magnitude, $m(1, 1, 0) = 15.35 \\pm 0.05$, measured when the object was active in late 2005 at a heliocentric distance of 2.58 AU, a $V$-band magnitude of $m_V(1, \\rh, 0) = 17.8$ is obtained \\citep{2009ApJ...694L.111H}.  Since the cometary activity in 176P is indicative of ice sublimation with a 30\\% contribution of the coma to the total brightness, the scaling relation between gas production rates and heliocentric magnitudes from \\citet{2008LPICo1405.8046J} predicts a water production rate of approximately $1.0 \\times 10^{26}\\ \\s$.  This correlation has been obtained for a sample of 37 comets with heliocentric distances between 0.32--4.53 AU.  Thus, if ice sublimation is the driving mechanism of 176P's activity, the derived \\ce{H2O} production rate is too low by about a factor of two to explain the activity level during its 2005 perihelion passage.  It is unlikely that sublimation of carbon monoxide ice is the source of the activity in MBCs since the temperature in the region where the comet formed could not have allowed condensation of CO. We note that the \\citet{2008LPICo1405.8046J} correlation should be taken with some care because none of the comet measurements were obtained at the brightness level of 176P.  From the dust production rate of 0.07 kg s$^{-1}$ estimated by \\citet{2011AJ....142...29H}, a water production rate of $2.3 \\times 10^{24}\\ \\s$ is inferred assuming a dust-to-gas ratio of one.  However, there are significant uncertainties in the value of the dust production rate determined from photometric data, and additionally the dependence on the dust-to-gas ratio with heliocentric distance is poorly constrained.  Mid-infrared photometric observations of 176P on 23--24 April 2010 by the {\\it Wide-field Infrared Survey Explorer} (WISE) mission did not provide any indication of a coma \\citep{2012ApJ...747...49B}, and there are no other published infrared or optical data closer to the last perihelion passage.  We conclude that water was not detected in our observation because the water production rate was lower than $\\upper \\times 10^{25}\\ \\s$ or the object was not active during our observation, so a more detailed study is needed to shed light on the activation mechanism in MBCs."
    },
    "1208/1208.4954_arXiv.txt": {
        "abstract": "We investigate the dynamical interaction of a central star cluster surrounding a super-massive black hole and a central accretion disk. The { dissipative force acting on stars in the disk} leads to an enhanced mass flow towards the super-massive black hole and to an asymmetry in the phase space distribution due to the rotating accretion disk. The accretion disk is considered as a stationary Keplerian rotating disk, which is vertically extended in order to employ a fully self-consistent treatment of stellar dynamics including the dissipative force originating from star-gas ram pressure effects. The stellar system is treated with a direct high-accuracy $N$-body integration code. A star-by-star representation, desirable in $N$-body simulations, cannot be extended to real particle numbers yet. Hence, we carefully discuss the scaling behavior of our model with regard to particle number and tidal accretion radius. The main idea is to find a family of models for which the ratio of two-body relaxation time and dissipation time (for kinetic energy of stellar orbits) is constant, which then allows us to extrapolate our results to real parameters of galactic nuclei. Our model is derived from basic physical principles and as such it provides insight into the role of physical processes in galactic nuclei, but it should be regarded as a first step towards more realistic and more comprehensive simulations. Nevertheless, the following conclusions appear to be robust: the star accretion rate onto the accretion disk and subsequently onto the super-massive black hole is enhanced by a significant factor compared to { purely stellar dynamical} systems neglecting the disk. This process leads to enhanced fueling of central disks in active galactic nuclei and to an enhanced rate of tidal stellar disruptions. Such disruptions may produce electromagnetic counterparts in form of observable X-ray flares. Our models improve predictions for their rates in quiescent galactic nuclei. We do not yet model direct stellar collisions in the gravitational potential well of the black hole, which could further enhance the growth rate of the black hole. Our models are relevant for quiescent galactic nuclei, because all our mass accretion rates would give rise to luminosities much smaller than the Eddington luminosity. To reach Eddington luminosities, outflows and feedback as in the most active QSO's other scenarios are needed, such as gas accretion after galaxy mergers. However, for AGNs close to the Eddington limit this process may not serve as the dominant accretion process due to the long timescale. ",
        "introduction": "Kinematic and photometric data of galactic nuclei have revealed that super-massive central black holes (SMBHs) are ubiquitous in most galaxies (see e.g. \\citealt{KormendyR:95} for a review). Detailed information on their photometric profiles and shapes of spectral line profiles allow us, in certain limits, to deduce the true shape of the distribution function in phase space of such systems. The distribution function determines the rate at which stars will come close to the SMBH, be tidally disrupted or destroyed by direct collisions and eventually accreted onto the black hole \\citep{FrankR:76,BahcallW:76}. Quasars, however, which are the most luminous witnesses of accretion activity onto SMBHs are observed already in the young universe at redshifts of $z>5$. It is difficult to explain how black holes can grow so quickly to the observed high masses ($10^6-10^9\\Msol$) by pure star accretion. It has been therefore argued that massive seed black holes form already by dissipative and viscous collapse, possibly accompanied by the formation of massive stars and their coalescence, at the time of galaxy formation \\citep{Colgate:67, SpitzerS:66, SpitzerS:67, Sanders:70, Rees:84}. In the hierarchical picture of galaxy formation, the most massive dark halos with small angular momentum can account for the early formation of the most massive black holes. The SMBH forms by a centrally focused collapse entering a phase of a dense super-massive gaseous object which is supported by star-gas interactions \\citep{BisnovatyiS:72,Vilkoviskij:75,Hara:78,Langbeinetal:90}. The interaction of a compact stellar cluster with a massive central object in the form of a superstar or SMBH was considered by \\citet{Vilkoviskij:75,Hills:75,Hara:78}. The evolution of the dense non-rotating stellar cluster was studied by \\citet{SpitzerS:66, Bisnovatyi:78} among others, and the evolution of the gas sphere was considered by \\citet{Langbeinetal:90}.  These studies have commonly neglected angular momentum, which should not be neglected during mergers nor in the intrinsic structure of galaxies. Spectra of active galactic nuclei suggest that there should be gas in the form of a massive central accretion disk (AD) in which all interstellar matter settles before it may be accreted. The origin of the AD could be inflowing cool gas during mergers and/or debris by direct stellar collisions or stellar evolutionary processes, depending on the evolutionary state of the host galaxy. \\citet{Artymowiczetal:93} provide a theoretical model framework in which star-gas interactions and the build-up of massive stars generate a massive AD around the central SMBH.  Recent work shows that gaseous disks in galactic centers around SMBHs are important for the dynamics and the morphology of central regions of galaxies and for the evolution of single or multiple central black holes in many ways. For example \\citet{Cuadra:09,Callegari:09,Callegari:10} study the role of small scale disks for the acceleration of binary black hole mergers, through dissipation of kinetic energy in the disk, after a galactic merger. \\citet{Baruteau:10} discuss the hardening of stellar binaries in circumnuclear disks and their subsequent interactions with central black holes, which may lead to high velocity stellar escapers.  \\citet{ShakuraS:73} developed a model for a stationary AD, which has been the basis of many investigations since then. Close to the inner boundary of the AD, at a few Schwarzschild radii $r_\\mathrm{S}$, general relativistic effects must be taken into account. \\citet{NovikovT:73} extended the standard disc model in that regime.  If the inner part of the AD reaches a critical surface density the interaction of orbits with the gas (or more correctly energy and momentum transfer of stars due to ram pressure effects, { henceforth denoted as dissipation}) cannot be neglected anymore. Stellar interactions with accretion disks were also considered by \\citet{VilkoviskijB:82, Vilkoviskij:83, Syeretal:91}. More detailed investigations of the stellar orbits, crossing accretion disks were presented in \\citet{VokrouhlickyK:98} and in \\citet{SubrK:99}. \\citet{SubrK:99} and \\citet{Subr:03} assumed an infinitely thin disk interacting with stars; they found that the interaction between the disk and the stars (star-gas { dissipation}) will deplete counter-rotating stars, create a flattening of the large-scale structure of the system and initiate anisotropies (i.e. changes of the eccentricity distributions). These studies did not take into account any feedback onto the disk or any finite-thickness effects of the disk. They tried to find a stationary state in which the transport of stars into the central galactic region is balanced by removal of stars from the central disk. Evolution timescales and initial conditions to reach this equilibrium was missing.  Around SMBHs the AD may extend to the parsec scale, orders of magnitude larger than $r_\\mathrm{S}$. \\citet{Rauch:95,Rauch:99} studied in detail the impact of the AD on the orbits and the distribution of stars with test particle simulations. Their work includes relativistic effects, even for the case of rotating Kerr black holes. They find that the orbital inclination with respect to the AD declines quickly as soon as the { dissipative} force becomes effective. Additionally they find a steepening of the central density cusp, due to the combined effect of relativistic orbit migration and stellar collisions.  The semi-analytic model by \\citet{VilkoviskiC:02} raises the point that two-body relaxation of the stars within and near the disk will tend to elevate trapped stars again out of the disk, and that the competition between relaxation and { dissipation} will define a stationary state of the system, with some well-defined stationary flux of stars going down to the black hole. \\citet{VilkoviskiC:02} compared the star-star two-body interactions with the star-disk interactions and concluded that the latter is stronger in the inner parts of the accretion disk and vice versa in the outer parts based on nearly circular orbits of the stars. They derived analytic approximations for the effective inclination of stellar orbits, where the inflow to the SMBH takes place. They also derived a critical radius inside which direct stellar collisions must be taken into account.  All of the mentioned papers so far have both strong and weak points. \\citet{Rauch:99} is the only paper to combine general relativity effects and stellar collisions, but considered only the central region where the (spherically symmetric) potential is dominated by the super-massive black hole. They used an approximate model of stellar dynamics based on a Monte Carlo technique, which requires a spherically symmetric central star cluster. In \\citet{Rauch:95} they combined relativistic effects with star-disk interactions, but the disk is assumed to be infinitely thin. The infinitely thin disk approximation was also used in other work by \\citet{SubrK:99,Subr:03,VilkoviskiC:02}).  Following the ideas of \\citet{VilkoviskiC:02} we investigate by self-consistent direct $N$-body simulations, the interaction of a central compact stellar cluster (CSC) with the AD in active galactic nuclei. We focus on the mutual interplay of two-body relaxation and the depletion of stars by the { dissipative} force in the AD as a secular long-term evolution of the stellar mass distribution and the velocity distribution function of the CSC. In this paper we report results on a new model of star-disk interactions in galactic nuclei. Our focus lies on the correct and accurate representation of the stellar orbital motion crossing the disk, by implementing a disk with its density distribution in a full three-dimensional $N$-body simulation. We have added the force and force time derivative to the standard Hermite scheme (see below) as a function of local disk density and velocity in three dimensions. This is the most general approach, and later it will allow us to include evolving models of the AD and appropriately model the mass and energy transfer between the AD and CSC.  Our first results concern the enhanced accretion rate onto the central SMBH due to the interaction with the AD, in the regime where the relaxation timescale is comparable to the dissipation timescale. This is, to our knowledge, the first approach to study the competition between relaxation and star-gas { dissipation} in a direct simulation. Since the direct simulations are not yet able to reach realistic particle numbers and spatial resolution, we will perform a careful scaling analysis to show in which way our numerical simulations have to be interpreted for real astrophysical galactic nuclei.  This paper is organized as follows: in Sect.~\\ref{sec-physics} we describe the accretion disc and the { dissipative} force in detail, Sect.~\\ref{sec-model} gives the numerical realizations of the system, in Sect.~\\ref{sec-result} the results are discussed and in Sect.~\\ref{sec-conclusion} a summary and conclusions are presented.  In follow-up papers we will discuss the dependence of the { dissipation} on the orbital parameters and the phase space evolution of the cusp stars in detail and take into account the feedback of the star-gas interaction on the AD properties. Detailed studies of migration of stars, binaries, and black holes inside the disk towards the center, and its observational consequences will also be included in future work. ",
        "conclusions": "\\label{sec-conclusion}  In galactic nuclei super-massive black holes coexist with a dense stellar cluster; galaxy mergers and the quasar phenomenon indicate that at least for some time there should be also large amounts of interstellar gas present in the nuclear regions around the black holes. In this paper we have examined the interaction and co-evolution of a dense star cluster surrounding a star-accreting super-massive black hole with an assumed central gaseous disk. Interactions of such disks with the surrounding dense star clusters have been proposed as source of gas supply to the central disk \\citep{MiraldaK:05}, as agents to enhance the tidal star accretion rate \\citep{VilkoviskiC:02}, and to cause feedback on the orbital parameters of stars \\citep{Rauch:95}, including a modification of sources of gravitational waves \\citep{Rauch:99}.  Our model is the first self-consistent long-term simulation of a dense star cluster, surrounding a star-accreting super-massive black hole, and subjecting the stars to the { dissipative} forces from a resolved central gaseous disk. We resolve effects of two-body relaxation, dissipation of stellar kinetic energy in the disk and star accretion onto the central black hole in a numerical study based on direct high-accuracy simulations. Since star-by-star modeling of a galactic nucleus down to the realistic tidal radius is not yet possible, despite of the modern GPU hardware used for simulations, a careful scaling analysis is presented as a function of the particle number in the simulations and of an assumed star accretion radius, to allow conclusions for the real astrophysical situation in galactic nuclei. But our model still has a number of serious drawbacks. Firstly, it assumes a stationary accretion disk, so energetic feedback to the disk structure by star-disk interactions is neglected, secondly, the physics of star-gas interactions is modeled approximately. Finally, there is no distinction between properties of different stars (main sequence, giants, remnants), but rather a single stellar species with solar radius is assumed. In that sense our study should be considered as a pathfinder and exploratory.  This investigation is a direct continuation of a semi-analytic study by \\citet{VilkoviskiC:02}, extending it by a more detailed and numerical study of the stellar dynamical evolution of the central stellar cluster. Our paper uses a numerical approach based on direct $N$-body simulations, including particle-particle forces as well as a { dissipative} force in the disk. Here we resolve the dissipation of stellar kinetic energy along the stellar paths in the vertically extended accretion disk. Particle numbers in our simulations and an adopted star accretion radius (onto the central super-massive black hole) are used as free parameters in our model, while for other important parameters of the problem fiducial values are assumed. These are e.g. the initial super-massive black hole mass (10 percent of the initial central stellar cluster mass), the gaseous accretion disk mass (10 percent of the initial black hole mass), and the outer radius of the accretion disk (set equal to the black hole gravitational influence radius). Finally, all star particles are equal in the simulation, and their total effective cross section is used as a parameter to obtain the physically correct ratio of dissipation to relaxation time.  We show that the accretion rate of stars onto the super-massive black hole is strongly enhanced by the { dissipative} force of the accretion disk (a factor of four with our parameters). The accretion process is determined by an equilibrium of diffusion by two-body encounters and energy loss by the { dissipative} force. Consistently there is also an energy deposition in the central accretion disk; we find that most stars accreted or trapped in the disk are quickly accreted also to the black hole, because there is no stable co-rotation of the stars with the disk. Our results are robust with respect to variation of particle number and adopted accretion radius, therefore they should hold under realistic conditions in galactic nuclei. Our star accretion rate does not depend strongly on the adopted star accretion radius; it supports the idea of \\citet{VilkoviskiC:02}, that there is a stationary flow of stars within the disk towards the central black hole, which is determined by an equilibrium between dissipation and relaxation time. In spite of the enhanced number of stars accreted through the disk, we still find that there is a Bahcall-Wolf central density cusp present in the system, which is not significantly perturbed.  Central densities in star clusters near super-massive black holes can reach $10^8 \\Msol\\,\\mathrm{pc}^{-3}$ or more. At such high stellar densities direct, disruptive stellar collisions may produce gas in the gravitational potential well of the black hole. The gas production rate could be larger than the one obtained from tidal disruption of stars \\citep{SpitzerS:66,SpitzerS:67}; see also \\citet{Begelman:78,FrankR:76}, and numerical models in e.g. \\citet{FreitagB:02,FreitagGR:06,FreitagRB:06}. But there is little doubt that a large fraction of this gas is finally accreted to the SMBH; some fraction of it though may escape. The same is true for the gas produced by tidal accretion. Our models do not yet resolve the very central regions of the star cluster, where stellar collisions may occur predominantly. We anyway treat the accretion radius as a free parameter used for scaling studies; it is here usually large compared to the astrophysically defined tidal radius, where stars are destroyed by tidal forces. Our results should not depend strongly on whether the stars inside $r_\\mathrm{acc}$ are finally disrupted by tidal forces or destroyed by stellar collisions with subsequent accretion of the gas onto the SMBH. Less is known about the effect of induced stellar collisions due to low relative velocities of stars in the disk (due to dissipation). We will study the effects of direct stellar collisions in future work.    Direct stellar collisions produce another source of gas deep in the gravitational well of the central supermassive black hole. As in the case of tidal disruptions of stars, we do not know exactly how much mass is accreted to the SMBH, and how much is ejected due to magnetic fields (jets) and radiation pressure near it \\citep[see e.g. one improved model by][]{Kasen:10}. Our assumption to add 100\\% of the material from tidal accretion to the black hole clearly is an upper limit, the real growth rate may be lower due to some mass loss in the process, also for stellar collisions. However, even in our case with possibly overestimated accretion rates, assuming a typical value of 10\\% mass to energy conversion, the luminosity obtained from our mass accretion rates are of the order of $10^8\\Lsol$, much smaller than the Eddington luminosity. Therefore our results are applicable to quiescent galactic nuclei, not to quasars or AGN in their most active phase, where high mass accretion rates, feedback and outflows are driven for example by gas inflow due to galaxy mergers.   The response of the central accretion disk to the deposition of energy and stars is neglected in this paper. We will work on this problem in future studies. The black hole growth rate due to star accretion will be limited by the condition that the accretion disk can find a new equilibrium absorbing the dissipated stellar energy and radiating it. In the inner spherically symmetric part of the system the Eddington limit will pose an upper limit to the star accretion rate allowed in a stationary state \\citep[see discussion by][]{MiraldaK:05}. In that sense our accretion rates, determined by neglecting feedback on the disk, will be upper limits. It is still possible that the process of star trapping to the disk and star and gas accretion to the central black hole is quasi-periodic and highly non-stationary (as suggested by the ubiquitous time variability of radiation from active galactic nuclei)."
    },
    "1208/1208.4486_arXiv.txt": {
        "abstract": "We present a detailed study of star formation occurring in bound star--forming clouds under the influence of internal ionizing feedback from massive stars across a spectrum of cloud properties. We infer which objects are triggered by comparing our feedback simulations with control simulations in which no feedback was present. We find feedback always results in a lower star--formation efficiency and usually but not always results in a larger number of stars or clusters. Cluster mass functions are not strongly affected by feedback, but stellar mass functions are biased towards lower masses. Ionization also affects the geometrical distribution of stars in ways that are robust against projection effects, but may make the stellar associations more or less subclustered depending on the background cloud environment. We observe a prominent pillar in one simulation which is the remains of an accretion flow feeding the central ionizing cluster of its host cloud and suggest that this may be a general formation mechanism for pillars such as those observed in M16. We find that the association of stars with structures in the gas such as shells or pillars is a good but by no means foolproof indication that those stars have been triggered and we conclude overall that it is very difficult to deduce which objects have been induced to form and which formed spontaneously simply from observing the system at a single time. ",
        "introduction": "The influence of feedback from stars on the star--formation process itself is a long--standing and intriguing problem. Such feedback is usually invoked in the negative sense of `self--regulating star formation' -- the disruption of giant molecular clouds (GMCs) and embedded clusters by massive stars and the consequent shutting down of star formation. However, stellar feedback also has a positive component in the sense of triggered star formation -- the inducement of GMCs by massive stars to form new stars that they would not otherwise give birth to (\\cite{1977ApJ...214..725E,1995ApJ...451..675E}, and see \\cite{2011EAS....51...45E} for a brief up--to--date review). It is highly likely that stellar feedback operates in both modes simultaneously and that it triggers the formation of additional stars in some regions of GMCs while expelling gas from others. The interesting question is then whether its \\textit{overall} effects are positive or negative.\\\\ \\indent Observations of triggered star formation are legion and are usually loosely divided according to two popular models. The collect--and--collapse model involves the fragmentation, via gravitational and other instabilities, of a shell of dense material swept up by an expanding feedback--driven bubble. If simplifying assumptions about the bubble geometry and the smoothness of the background gas are made, this process lends itself easily to analytical \\citep[e.g.][]{1994A&A...290..421W,2001A&A...374..746W} and numerical \\citep[e.g.][]{2007MNRAS.375.1291D} study. There is also a large and growing body of observational work on this topic \\citep[e.g.][]{2006A&A...446..171Z,2006A&A...458..191D,2008A&A...482..585D,2010A&A...518L..81Z} with which to compare the theoretical work. The case for triggered star formation in these systems is compelling. However, recent work by \\cite{2011arXiv1109.3478W} questions whether the existence of a smooth shell is necessarily a pointer to the collect and collapse process in action. They perform simulations of HII regions expanding into fractal molecular clouds with various fractal dimensions and find that shell--like structures can readily be produced in even quite strongly non--uniform background clouds and reflect the initial gas distribution, not the collect and collapse process.\\\\ \\indent If the assumptions of relatively simple (usually spherical) geometry and homogeneous ambient gas are dropped, as they must be in turbulent and highly non--uniform GMCs, identification of triggered stars becomes rather more difficult. Expanding HII regions (and wind/supernova bubbles) then encounter pre--existing structures in the surrounding gas, which may or may not be gravitationally unstable already and which will either be destroyed or induced to collapse. The latter outcome is described by the radiation--driven--implosion model (studied intensively by, e.g. \\cite{1994A&A...289..559L,2003MNRAS.338..545K,2009MNRAS.393...21G,2011ApJ...736..142B}), although a stellar wind or supernova shock may have the same eventual result \\citep[e.g.][]{1996ApJ...468..784F,1998ApJ...508..291V}. Disentangling triggered from spontaneous star formation (i.e. star formation that was going on anyway) in these circumstances becomes very difficult. The association of young stellar objects (YSOs) with shells/cavities \\citep[e.g.][]{2003ApJ...595..900K,2008ApJ...688.1142K,2009A&A...503..107P} or pillars \\citep[e.g.][]{1999AJ....117..225W,2005AJ....129..888S,2007ApJ...654..347L}, their proximity to ionization fronts \\citep[e.g.][]{2009ApJ...700..506S} and the existence of bright--rimmed clouds \\citep[e.g.][]{1995ApJ...455L..39S,2009A&A...497..789U} have all been used to infer star formation induced by stellar feedback but in all of these objects, the gas morphology and distribution of stars are very complex and difficult to interpret.\\\\ \\indent Triggering may also be inferred more generally by searching for instances of sequential or self--propagating star formation, which can in principle be inferred by looking for spatial age gradients in star--forming regions or complexes. This idea was first proposed by \\cite{1978SvAL....4...66E} and was confirmed on large scales (hundreds of pc to $\\sim$1 kpc) by \\cite{1989SvAL...15..388S}. On the scale of single associations, \\cite{1995ApJ...451..675E} did the seminal theoretical work and considerable observational work followed. \\citep[e.g.][]{1985ApJ...297..599D,2005AJ....129..776N,2006PASJ...58L..29M,2010ApJ...713..883B}. Recently, attempts have been made to look for statistical correlations between the positions of young stellar objects and infra--red bubbles. \\cite{2012MNRAS.421..408T} and \\cite{2012arXiv1203.5486K} both find statistically--significant overdensities of YSOs within and especially on the borders of, young feedback--driven bubbles. These groups infer that a few tens of percent of all massive stars in the Milky Way may have been triggered.\\\\ \\indent In two previous papers, \\cite{2007MNRAS.375.1291D}, \\cite{2012MNRAS.tmp.2723D}, we have attempted to contribute to this discussion by modelling the effect on star formation in (unbound and bound, respectively) turbulent GMCs of external radiation by an arbitrarily--placed O--star. By comparing with control runs in which feedback was absent, we were able to quantify the degree of triggering, which we found to be modest (increasing the star--formation efficiency by $\\sim30\\%$ at most) in both cases. In this paper, we seek to model triggering in a more realistic setting and study the influence on the star formation in GMCs by stars that have already formed within that cloud. We take as our starting point a series of simulations described in \\cite{2012MNRAS.424..377D}, hereafter Paper 1, whose purpose was to assess how efficient O--star photoionization alone could be in disrupting bound GMCs.\\\\ \\indent In Paper 1, we describe in some detail the overall dynamical reaction of our model clusters to the ionizing feedback of the O--stars or O--star--hosting clusters formed within. Models C, G and H form no stars at all. Models E, F, B and X form stars/clusters vigorously but a combination of dense gas and strong accretion flows stifling the ionization of fresh gas, and the large escape velocities of these systems reduced the impact of ionization severely. In contrast, runs A, D, I and J were strongly affected by feedback, with several tens of percent of their gas reserves being expelled in the canonical 3Myr time window before the first supernova explosions. However, we noted that in the none of the clouds was feedback able to bring star formation to a halt, although star formation was noticeably slowed in the runs where feedback had tangible effects. However, these simulations also exhibited many morphological features that are often taken to be signposts of triggered star formation. We suggested in Paper 1 that the negative impact of gas expulsion might be to some extent counterbalanced by triggering of star formation -- i.e. by the birth of stars/clusters in the feedback simulations which would not have formed in the absence of feedback. In this paper, we investigate this possibility in detail by performing control runs identical to runs A, D, I and J in all respects except that photoionization was forbidden. We find that triggered star/cluster formation occurs in all our simulations, but that the overall effect of feedback on the star--formation efficiency is always negative.. ",
        "conclusions": "\\indent We have investigated the effects of ionizing feedback from massive stars on the star--formation process in four model molecular clouds known from Paper I to be significantly influenced by feedback. We find that, although there is strong evidence of triggered star formation and even of occasional triggered cluster formation, the overall effect on the star--formation efficiency is always to decrease it. The expulsion of potentially--star forming gas and the disruption of accretion flows feeding objects that have already formed always outweighs triggering in these calculations.\\\\ \\indent Feedback of this form is able to significantly influence the stellar IMF, either by simply shifting it to lower masses by curtailing accretion onto all stars, as mostly seen in Run I,  or by forming an excess of new stars and preventing them from accreting, as seen in Run J. Feedback also limits the growth of the most massive stars. However, the IMFs so formed are scarcely unusual in appearance and there is no obvious way of inferring the presence or degree of triggering purely from observing the mass function in a given system.\\\\ \\indent Similarly, although we find that ionizing feedback can profoundly alter the geometrical distribution of the stars within a cluster in ways that are robust against projection effects, the degree to which, and the direction in which, it does so depend on the cluster environment. It is not the case that internally--triggered star formation always leads to more subclustering (as measured quantitatively by the Q parameter), since it may produce a relatively smooth halo of triggered stars around an otherwise structured central cluster. In Run I, the density is low so that it takes a long time for the expanding bubbles to sweep up shells which are gravitationally unstable -- most of the triggering in Run I on the few Myr timescales considered here therefore happens in the accretion flows which dominate the cloud's structure. In Run J, the background density is higher and the sweeping up of material by expanding shells efficiently produces widely distributed triggered stars. Feedback may thus increase or decrease the Q parameter depending on the properties of the background cloud.\\\\ \\indent Despite the difficulty of inferring triggering from the geometry of the stars alone, the \\emph{combined} geometry of the stars and the surviving cold gas provides an indication of where induced star formation is occurring, by the association of triggered stars with shells or pillars, but even this correlation is not always reliable. We suggest that the partial destruction of accretion flows feeding dense gas towards ionizing sources may be a natural explanation for pillar--like structures.\\\\ \\indent Triggering of star formation by ionization is seen in our simulations, but does not compensate for the decrease in star formation due to gas expulsion by the same process. Discerning which stars are triggered remains problematic and their overall numbers are only marginally significant compared to those generated by ongoing spontaneous star formation. In particular, we find that it is not not always safe to infer that individual stars have been triggered merely from their association with pillars or bubble walls."
    },
    "1208/1208.1606_arXiv.txt": {
        "abstract": "The light deflection of one component of a binary system due to the gravitational field of the other component is investigated. While this relativistic effect has not been observed thus far, the question arises that whether this effect becomes detectable in view of todays high-precision astrometry which soon will reach the microarcsecond level of accuracy. The effect is studied and its observability is investigated. It turns out, that in total there are about $10^3$ binaries having orbital parameters such that the light deflection amounts to be at least $1$ microarcsecond. Two stringent criteria for the orbital parameters are presented, by means of which one can easily determine the maximal value of light deflection effect for a given binary system. It is found, that for relevant binaries their orbital parameters must take rather extreme values in order to have a light deflection of the order of a few microarcseconds. Only in a very few and rather extreme binary systems the light deflection effect might be detectable by todays astrometry, but their existence is highly improbable. Thus, the detection of this subtle effect of relativity still remains a challenge for future astrometric missions. ",
        "introduction": "\\label{Section0}  Astrometric space missions, especially the ESA (European Space Agency) cornerstone mission Gaia, see e.g. {\\it Perryman, et al.} \\cite{Gaia_Overview}, are in preparation to attain microarcsecond ($\\mu{\\rm as}$) level of accuracy in absolute positional measurements of stars and other celestial objects. This unprecedented accuracy of astrometric observations makes it necessary to account for many subtle effects which were totally negligible before. A practical model for astrometric observations with an accuracy of $1\\,\\mu{\\rm as}$ has been formulated by {\\it Klioner} \\cite{Klioner1}, where the influence of the gravitational fields inside the solar system were taken into account. Furthermore, a number of additional effects potentially observable at this level of accuracy due to various gravitational fields generated outside of the solar system were also briefly discussed in that investigation. One of them is the gravitational light deflection of one companion of a binary system in the gravitational field of the other companion (without loss of generality, throughout the investigation the light deflection of component B at component A is considered, hence, component A is considered to be the massive body, while component B is the light-source). This effect would change the apparent position of one component of a binary system. While it is clear that this effect is relatively small and even at the level of $1\\,\\mu{\\rm as}$ observable only for edge-on binary systems, the huge amount of binaries generate some hope that there are relevant systems where this light deflection effect becomes detectable. For instance, Gaia will observe $10^9$ stars brighter than $20^{\\rm th}$ apparent magnitude. Detailed numerical simulations predict the detection of about $10^8$ (resolved, astrometric, eclipsing, spectroscopic) binary systems by Gaia mission, see {\\it Zwitter {\\rm \\&} Munari} \\cite{GAIA_Binary1}, which is a considerable increase compared to the $10^5$ binary systems known so far, see \"{\\it Washington Double Star Catalog}\" \\cite{Washington_Double_Star}.  Let us consider this argument a bit more quantitatively. For a simple estimate of the expected order of magnitude in light deflection, the classical lens equation (in the form given by Eq.(67) in {\\it Fritelli, et al.} \\cite{Fritelli_Kling_Newmann}, Eq.~(24) in {\\it Bozza} \\cite{Bozza} or Eq.~(23) in {\\it Zschocke} \\cite{Article_Generalized_Lens_Equation}) is applied. In terms of orbital elements of a binary system it can be written as follows: \\begin{eqnarray} \\varphi = \\frac{1}{2} \\left(\\sqrt{\\frac{A^2}{r^2}\\,\\cos^2 i + 16\\,\\frac{m}{r}\\,\\frac{A}{r}\\,\\sin i} - \\frac{A}{r}\\,\\left|\\cos i\\right|\\right). \\label{classical_lens_1} \\end{eqnarray}  \\noindent Here, $\\varphi$ is the light deflectin angle, $i$ is the inclination, $A$ is the semi-major axis and $r$ is the distance of center-of-mass of the binary system from the observer, and the Schwarzschild radius is $\\displaystyle m= \\frac{G\\,M}{c^2}$, where $G$ is the gravitational constant and $c$ is the speed of light, and $M$ is the stellar mass of component A of the binary system. From (\\ref{classical_lens_1}) one obtains the maximal possible light deflection for edge-on binaries, i.e. attained for the case when the inclination is exactly $90^{\\circ}$: \\begin{eqnarray} \\varphi &\\le& 200\\,\\mu{\\rm as}\\;\\sqrt{\\frac{M}{M_{\\odot}}\\;\\frac{A}{\\rm AU}}\\;\\frac{\\rm pc}{r}\\;. \\label{classical_lens_2} \\end{eqnarray}  \\noindent Here, ${\\rm AU} = 1.496 \\times 10^{11}\\;{\\rm m}$ is the astronomical unit, ${\\rm pc} = 3.086 \\times 10^{16}\\;{\\rm m}$ stands for parallax of one arcsecond, and $M_{\\odot}$ is the solar mass. According to this formula, the choice of moderate values like $M=M_{\\odot}$ and $A \\sim 100\\,{\\rm AU}$ would result into significant light deflection effect on microarcsecond level even at large distances of about $r \\sim 100 \\,{\\rm pc}$. A meaningful value for the density of stars in the solar neighborhood, $0.025 \\,{\\rm binaries}\\;{\\rm pc}^{-3}$, implies already $10^5$ binaries inside a sphere of $r = 100\\,{\\rm pc}$. Thus, one might conclude among them there are a few relevant edge-on binary systems.  However, in order to estimate quantitatively the number of relevant systems, one needs to know the probability for the occurrence of such edge-on binaries which have a given light deflection depending on their orbital parameters like inclination, mass, distance and semi-major axis. Such a relation between a given light deflection and orbital parameters is given by a so-called inclination formula. Since the inclinations of binary systems are of course randomly distributed, it is meaningful to resolve such an inclination formula in terms of inclination. As it has been shown by {\\it Klioner, et al.} \\cite{inclination_formula} (see Appendix \\ref{Appendix_KMS} for some basic steps), such an inclination formula can be obtained by means of the analytical solution of light deflection in standard post-Newtonian approach, and is given as follows: \\begin{eqnarray} \\left|\\frac{\\pi}{2} - i \\right|_{\\rm KMS} &\\le& 2\\,\\arctan \\left(0.0197\\;\\frac{M}{M_{\\odot}}\\; \\frac{\\mu{\\rm as}}{\\varphi}\\;\\frac{\\rm pc}{r}\\right). \\label{KMS_D} \\end{eqnarray}  \\noindent For a better illustration of the inclination formula, relation (\\ref{KMS_D}) is rewritten in terms of angular degrees instead of radians: \\begin{eqnarray} \\left| \\,90^{\\circ} - i\\,\\right| &\\le& 2.25^{\\circ}\\;\\frac{M}{M_{\\odot}}\\;\\frac{\\mu{\\rm as}}{\\varphi}\\;\\frac{\\rm pc}{r}\\;, \\label{condition_0} \\end{eqnarray}  \\noindent where also $\\arctan x = x + {\\cal O} (x^3)$ has been used. According to this relation, the inclination $i$ of a binary system with stellar mass $M$ and at distance $r$ must not deviate from the edge-on value $90^{\\circ}$ too much in order to have a given light deflection $\\varphi$. For example, for a hypothetical binary star with $M=M_{\\odot}$ situated at a distance of $r=10\\,{\\rm pc}$ the light deflection effect attains $1\\,\\mu{\\rm as}$ only if $\\left| \\,90^{\\circ} - i\\,\\right| < 0.225^{\\circ}$, which means that the probability to observe this binary at a favorable inclination is only about $0.2\\%$.  But even this estimation is still much too optimistic. As a concrete example of todays high-precision astrometry, let us consider one important parameter about the astrometric accuracy of the Gaia mission: the accuracy of one individual positional measurement in the most ideal case (bright star, i.e. $10^{\\rm th}$ magnitude) amounts to be $25\\,\\mu{\\rm as}$, which implies $\\varphi \\ge 25\\,\\mu{\\rm as}$. Furthermore, inside a sphere of $10\\,{\\rm pc}$ around the Sun almost every star and binary system is known already by the data of {\\it Research Consortium on Nearby Stars (RECONS)} \\cite{Recons}. Since inside that sphere there is no binary system having a light deflection on microarcsecond level, one has to take at least $r > 10\\,{\\rm pc}$. By taking into account these both remarks, one obtains $\\displaystyle \\frac{\\mu{\\rm as}}{\\varphi}\\;\\frac{\\rm pc}{r} = \\frac{1}{250}$ in relation (\\ref{condition_0}). Therefore, even in the best case one has to conclude $\\left| \\,90^{\\circ} - i\\,\\right| \\le 0.01^{\\circ}\\;M/M_{\\odot}$, that means (for $M=M_{\\odot}$), the probability to observe such a binary at a favorable inclination is practically only about $0.01\\%$. And the binaries must be, in fact, almost edge-on in order to have a light deflection which can be detected by todays astrometry. Accordingly, while relation (\\ref{classical_lens_2}) triggers the hope about the existence of many relevant binary systems, from relation (\\ref{condition_0}) one concludes that the number of relevant binary systems is considerably reduced.  An estimation of the total amount of binaries depends on many different parameters, like mass, semi-major axis, inclination and distances of the binaries. Therefore, a simple estimation is not so straightforward as one might believe. Moreover, the relation (\\ref{classical_lens_2}) has been obtained with the aid of classical lens equation, while relation (\\ref{condition_0}) has been obtain by means of standard post-Newtonian approach. These both approaches have different regions of validity. However, a rigorous treatment of the problem of light deflection in binary systems implies the need of an analytical formula which is valid for such kind of extreme astrometric configurations like the binary systems are. Recently, a generalized lens equation has been derived by {\\it Zschocke} \\cite{Article_Generalized_Lens_Equation}, which allows to determine the light deflection of binary systems on microarcsecond level. One aim of this study is, therefore, to reobtain the criteria (\\ref{classical_lens_2}) and (\\ref{condition_0}) as stringent conditions from one and the same approach, i.e. with the aid of generalized lens equation. This is possible, because in the corresponding limits the generalized lens equation agrees with the classical lens equation and the standard post-Newtonian solution. Another aim is, to derive an inclination formula like (\\ref{KMS_D}) for binary systems which allows to determine the needed inclination for a given light deflection angle and as a function of the orbital parameters. For that one has to take into account the distribution of stellar masses and the distribution of semi-major axes in binary systems. Finally, the aim of this study is, to determine the total number of relevant binaries having a light deflection on microarcsecond level, and to investigate the possibility to detect this effect of light deflection by todays high precision astrometry.  The article is organized as follows: In Section \\ref{Section1} some basics about orbital elements of binary systems are given. The generalized lens equation and the inclination formula are presented in Section \\ref{Section2}. In Section \\ref{Condition1} two stringent conditions on the orbital elements of binary systems (astrometric, spectroscopic, eclipsing and resolved binaries) are presented, which allow to determine whether or not the binary system will have a light deflection of a given magnitude. The total number of binaries which have a given light deflection for an infinite time of observation is estimated in Section \\ref{Section3a}, while the more practical case of a finite time of observation is considered in Section \\ref{Section3b}. The special case of resolved binaries is considered in Section \\ref{Condition2}. For that, the specific instrumentation of Gaia mission is considered in some detail as the most modern astrometric mission with todays highest possible accuracy. A Summary is given in Section \\ref{Summary}. ",
        "conclusions": "\\label{Summary}  In this study, the light deflection in binary systems has been considered. While there is absolutely not any doubt about the existence of this relativistic effect, is has not been observed so far. To investigate this effect of light deflection, an inclination formula (\\ref{inclination_30_A}) has been derived by means of generalized lens equation (\\ref{generalized_lens_1}) obtained recently by {\\it Zschocke} \\cite{Article_Generalized_Lens_Equation}, and these both equations are the theoretical basis for investigating the light deflection effect in binary systems. A simplified inclination formula has been presented by Eq.~(\\ref{inclination_50}) and its validity has been discussed in some detail. This simplified inclination formula has also been obtained by {\\it Klioner, et al.} \\cite{inclination_formula} by an independent approach. Furthermore, two stringent conditions on the orbital parameters have been given by Eqs.~(\\ref{condition_1}) and (\\ref{condition_2}). These both stringent conditions are relations between the orbital elements of a (resolved, astrometric, eclipsing, spectroscopic) binary system for a given magnitude of light deflection, and allow to find a relevant binary system in a straightforward way.  In Section \\ref{Section3a}, the total number of binaries with a given light deflection has been determined by means of the semi-major axis distribution according to \"{\\it \\\"Opik's law}\" and the mass distribution according to \"{\\it Salpeter's mass distribution}\". Since the inclinations are randomly distributed, the inclination formula allows to estimate the total number of relevant binaries with the aid of Eq.~(\\ref{numerical_results_5}). It turns out, that in total there exist about $10^3$ binaries having orbital parameters such that the light deflection amounts to be at least $1\\,\\mu{\\rm as}$; see FIG.~\\ref{FIG: Number_Binaries_2}.  In Section \\ref{Section3b} a finite time of observation of $5$ years (Gaia mission time) has been considered, which considerably reduces the total number of relevant binaries. Clearly, this case is of more practical importance, since a restricted time window of observation is in more agreement with reality than the first scenario. By taking into account the probability to find the system in the ideal astrometric position $E=0$ where the light deflection becomes maximal, it has been found by evaluating the corresponding integral (\\ref{observable_5}) that there is not any of the relevant binary systems in the ideal position $E=0$ during Gaia mission time; see FIG.~\\ref{FIG: Number_Binaries_3}. Thus, while in principle a few binaries will have a significant light deflection, the effect could not be detected due to the restricted finite time window of observation.  Furthermore, the special case of resolved binaries has been considered in Section \\ref{Condition2}. The astrometric instrumentation of the ESA cornerstone mission Gaia, see e.g. {\\it Perryman, et al.} \\cite{Gaia_Overview}, has been considered in some detail in order to decide whether or not this subtle effect of light deflection can be observed. Two conditions for resolved binaries were presented in Eqs.~(\\ref{condition_3}) and (\\ref{condition_4}) for such special kind of binary systems. It has been shown, however, that even for the Gaia mission, which is an outstanding milestone of progress in astrometry, such binary systems must have rather extreme orbital parameters in order to reach todays level of detectability, i.e. on microarcsecond level. The existence of such exotic binaries is, however, highly improbable.  In summary, the main results are presented by the inclination formulae (\\ref{inclination_30_A}) and its simplified version (\\ref{inclination_50}), the stringent conditions (\\ref{condition_1}) and (\\ref{condition_2}) and by the diagrams FIG.~\\ref{FIG: Number_Binaries_2} and FIG.~\\ref{FIG: Number_Binaries_3}. Accordingly, one comes to the conclusion that the detectability of light deflection in binary systems reaches the technical limit of todays high precision astrometry and might be detected only in case of a very few and highly exotic binary systems. It is, however, very unlikely that such extreme binaries might exist. It seems that the detection of the light deflection effect in binary systems needs an astrometric accuracy of better than about $0.1 \\mu{\\rm as}$. Thus, only astrometric missions of the next generation can accept the challenge to detect this subtle effect of relativity."
    },
    "1208/1208.0830_arXiv.txt": {
        "abstract": "Ground-based telescopes equipped with adaptive-optics (AO) systems and specialized science cameras are now capable of directly detecting extrasolar planets. We present the expected scientific capabilities of CHARIS, the Coronagraphic High Angular Resolution Imaging Spectrograph, which is being built for the Subaru 8.2 m telescope of the National Astronomical Observatory of Japan. CHARIS will be implemented behind the new extreme adaptive optics system at Subaru, SCExAO, and the existing 188-actuator system AO188. CHARIS will offer three observing modes over near-infrared wavelengths from 0.9 to 2.4 $\\mu$m (the $y$-, $J$-, $H$-, and $K$-bands), including a low-spectral-resolution mode covering this entire wavelength range and a high-resolution mode within a single band. With these capabilities, CHARIS will offer exceptional sensitivity for discovering giant exoplanets, and will enable detailed characterization of their atmospheres. CHARIS, the only planned high-contrast integral field spectrograph on an 8m-class telescope in the Northern Hemisphere, will complement the similar instruments such as Project 1640 at Palomar, and GPI and SPHERE in Chile. ",
        "introduction": "\\label{sec:intro} The past seventeen years have seen the discovery of over 700 planets around stars other than the Sun, commonly called ``exoplanets'' (see \\verb|www.exoplanets.org|).  Most of these exoplanets have been discovered by the radial velocity variations they induce in their host stars, or through periodic stellar dimming due to planetary transits. Both methods are most sensitive to planets on short-period orbits.  The magnitude of a radial velocity signal decreases with increasing orbital period, while large-separation planets must be exquisitely aligned to transit our line-of-sight from Earth.  In addition, both of these indirect methods generally require follow-up for one or more orbital periods, which could be centuries for a long-period exoplanet.  With the development of high-contrast, high angular resolution imaging, large telescopes can now directly image giant planets at large distances from their host stars \\cite{marois_et_al2008,kalas_et_al2008,lagrange_et_al2009,kraus+ireland2011}. The addition of spectroscopy to direct imaging will enable the characterization of giant exoplanet atmospheres and the development of techniques to find Earth-like planets that may support life\\cite{kawahara_et_al2012}.  This contribution describes the science case for CHARIS, a new Integral Field Spectrograph (IFS) designed for taking spectra of exoplanets, to be built for the Subaru 8.2 m telescope.  CHARIS will simultaneously obtain spatial and spectral information over the field-of-view (FOV) by dispersing the entire image on the detector. CHARIS will be the first high-contrast IFS on an 8m-class telescope in the Northern Hemisphere, and will achieve an inner-working angle $\\sim$$2 \\lambda/D$ and contrasts of up to $10^{-7}$.  CHARIS will provide both a low-resolution (${\\rm R}=14$) mode in which it will simultaneously collect photons from $\\lambda$=0.9-2.4 $\\mu$m, and a high resolution (${\\rm R}=65$) mode over a single near-infrared bandpass: $y$-, $J$-, $H$- or $K$-band.  The technical design of the instrument is described in the accompanying paper\\cite{peters_et_al2012}. ",
        "conclusions": "CHARIS will be the first IFS for exoplanet studies on an 8m class telescope in the Northern Hemisphere.  It will achieve a small inner working angle ($2 \\lambda/D$), and high contrasts of up to $10^{-7}$, representing a factor of $\\gtrsim$10 improvement over the current, second-generation high-contrast instruments on the Subaru telescope. CHARIS will provide ${\\rm R} = 33$ spectral resolution over a $1.\\!\\!''75 \\times 1.\\!\\!''75$ FOV. It will also provide both a low-spectral-resolution mode (${\\rm R}=14$), able to collect imagery from $0.9-2.4$ $\\mu m$, and a high-spectral-resolution mode (${\\rm R}=65$) over a single near-infrared bandpass.  CHARIS will offer exceptional sensitivity to detect new exoplanets, and high contrast and spectral resolution to provide detailed characterizations of their atmospheres.  It represents a major advancement in exoplanet science and will help address uncertainties in the frequency, properties, and formation mechanism of giant exoplanets.  The instrument should achieve first light at the Subaru telescope by the end of 2015."
    },
    "1208/1208.0237_arXiv.txt": {
        "abstract": "{The close-in planet orbiting \\gl\\ presents a puzzling orbital eccentricity ($e\\simeq 0.14$) considering its very short orbital period. Given the age of the system, this planet should have been tidally circularized a long time ago. Many attempts to explain this were proposed in recent years, either involving abnormally weak tides, or the perturbing action of a distant companion.}{In this paper, we address the latter issue based on Kozai migration. We propose that \\glb\\ was formerly located further away from the star and that it underwent a migration induced by a massive, inclined perturber via Kozai mechanism. In this context, the perturbations by the companion trigger high amplitude variations to \\glb\\ that cause tides to act at periastron. Then the orbit tidally shrinks to reach its present day location.} {We numerically integrate the 3-body system including tides and General Relativity correction. We use a modified symplectic integrator as well a fully averaged integrator. The former is slower but accurate to any order in semi-major axis ratio, while the latter is first truncated to some order (4$^\\mathrm{th}$) in semi-major axis ratio before averaging.} {We first show that starting from the present-day location of \\glb\\ inevitably leads to damping the Kozai oscillations and to rapidly circularizing the planet. Conversely, starting from 5-10 times further away allows the onset of Kozai cycles. The tides act in peak eccentricity phases and reduce the semi-major axis of the planet. The net result is a two fold evolution, characterized by two phases: a first one with Kozai cycles and a slowly shrinking semi-major axis, and a second one once the planet gets out of the Kozai resonance characterized by a more rapid decrease. The timescale of this process appears in most cases much longer than the standard circularization time of the planet by a factor larger than 50.} {This model can provide a solution to the eccentricity paradox of \\glb. Depending on the various orbital configurations (mass and location o the perturber, mutual inclination\\ldots), it can take several Gyrs to \\glb\\ to achieve a full orbital decrease and circularization. According to this scenario, we could be witnessing today the second phase of the scenario where the semi-major axis is already reduced while the eccentricity is still significant. We then explore the parameter space and derive in which conditions this model can be realistic given the age of the system. This yields constraints on the characteristics of the putative companion.} ",
        "introduction": "The M-dwarf GJ 436 has been the subject of growing interest in recent years. This star is known to host a close-in Neptune-mass planet \\citep[Gl 436b][]{but04} that was furthermore observed to undergo transit \\citep{gil07}.  The monitoring of the transits of \\glb\\ helped constraining its orbital solution. Noticeably, it appears to have a significant non-zero eccentricity: $e=0.14\\pm0.1$ \\citep{man07}, furthermore refined to $e=0.14\\pm0.01$ \\citep{dem07}. This eccentricity is abnormally high for a small period planet \\citep[$P\\simeq 2.64\\,$days ][]{bal10}.  With such a small orbital period, tidal forces are expected to circularize the orbit within much less than the present age of the system \\citep[1--10~Gyr][]{tor07}. Tidal forces seem indeed to be at work in this system. The detection of the secondary transit of the planet enabled \\citet{dem07} to derive a brightness temperature of $T=712\\pm36\\,$K, with flux reradiated across the day-side only. Conversely, a stellar irradiation / thermal re-radiation balance leads for to an equilibrium temperature for the planet $T_\\mathrm{eq}=642\\,$K, assuming $T_\\mathrm{eff}=3350$\\,K for \\gl\\ and zero albedo for the planet. According to \\citet{demi07}, the temperature difference indicates tidal heading of \\glb, but this could alternatively be due to the $8\\,\\mu$m sampling the atmosphere in a hot band pass, if the planetary atmosphere does not radiate like a black-body.  Many theories were proposed in recent years to account for the residual eccentricity of \\glb. The most straightforward one is that the tides are sufficiently weak and/or the age of the system is small enough to allow a regular tidal circularization not to be achieved \\emph{yet}. This idea was proposed by \\citet{mar08} for \\glb, after \\citet{tri00} had suggested that this accounts for any close-in exoplanet observed with significant eccentricity. It is in fact related to our poor knowledge of the planet's tidal dissipation $Q_p$. $Q_p$ is a dimensionless parameter related to the rate of energy dissipated per orbital period by tidal forced oscillations \\citep{bar09}. The smaller $Q_p$, the more efficient the tidal dissipation is. To lowest order, the circularization time-scale $t_\\mathrm{circ}$ of an exoplanet reads \\citep{gol66,jack08} \\begin{equation} t_\\mathrm{circ}=\\frac{4}{63}\\frac{a^{13/2}}{\\sqrt{GM_*^3}}{Q_pm_p}{R_p^5}\\qquad, \\label{tcirc} \\end{equation} where $M_*$ is the stellar mass, $a$ is the planet's orbital semi-major axis, $m_p$ its mass and $R_p$ its radius. Assuming $Q_p=10^5$, $R=27,600\\,$km, $a=0.0287\\,$AU, $m=23.2\\,M_\\oplus$ and $M=0.452\\,M_\\odot$ \\citep{mar08}, this formula gives $t_\\mathrm{circ}=4.7\\times10^7\\,$yr, which is obviously less than the age of the system. But $Q_p$ is very badly constrained. For Neptune, it is estimated between $\\sim 10^4$ and $3.3\\times10^5$ \\citep{ban92,zha08}. We can thus consider $10^5$ as a standard likely value for \\glb, but \\citet{mar08} argues that $Q_p$ could be as high as a few $10^6$. In this context, assuming the lower bound (1\\,Gyr) for the age of the system, the circularization could not be achieved yet. This depends however on the starting eccentricity at time zero. \\citet{wis08} performed analytical calculations of energy dissipation rates for synchronized bodies at arbitrary eccentricities and obliquities. His concluded that for a large enough starting eccentricity ($\\ga 0.4$, as will be the case below), the circularization time-scale should be significantly reduced with respect to Eq.~(\\ref{tcirc}). In this context, the conclusions by \\citet{mar08} may no longer hold. This shows also that conclusions drawn on asymptotic formulas like Eq.~(\\ref{tcirc}) must be treated with caution. Numerical work is required.  Alternatively, many authors proposed that the eccentricity of \\glb\\ is sustained by perturbations by an outer massive planet, despite a standard $Q_p$ value. Such a long period companion was initially suspected as \\citet{man07}'s radial velocity data revealed a linear drift. More data did not confirm the trend, which is now believed to be spurious. \\citet{mon09} gives instead observational uppers limits for the putative companion that rule out any companion in the Jupiter mass range up to a few AUs.  Perturbers can be either resonant or non-resonant. Looking for a for a planet locked in mean-motion resonance with \\glb\\ appears a promising idea, as resonant configuration usually trigger larger eccentricity modulations. This was suggested by \\cite{rib08}, who fitted in the residuals of the radial velocity data a $\\sim5\\,M_\\oplus$ planet locked in 2:1 mean motion resonance with \\glb. However, this additional planet was invalidated furthermore. Dynamical calculations by \\citet{mar08} showed that this planet cannot sustain the eccentricity of \\glb\\ unless $Q_p\\ga10^6$. Moreover, \\citet{alon08} showed that a $\\sim 5\\,M_\\oplus$ planet in 2:1 resonance with \\glb\\ should trigger transit time variations that should have been detected yet. Such variations are unseen today with a high level of confidence \\citep{pont09}. Actually the constraints deduced from radial velocity monitoring, transit time and geometry monitoring \\citep{bal10,mon09} almost rule out an additional large planet locked in a low-order mean-motion resonance with \\glb\\ ($P<8.5\\,$days), and to have a Jovian like planet up to a few AUs.  Conversely, \\citet{tong09} investigated the effects of secular planetary perturbations in eccentricity pumping, considering both non-resonant (secular) and resonant configurations. They found some perturbers configurations that could account for the present eccentricity of \\glb\\ while still being compatible with the observational constraints, but with no tides. Incorporating the tides leads inevitably to damp all the eccentricity modulation and to circularize the orbit. Meanwhile, \\citet{bat09} reanalyzed carefully this issue. They found that as expected, in most cases the eccentricity of \\glb\\ quickly falls to zero thanks to tidal friction, but that this effect can be considerably delayed if the 2 planet system lies initially in a specific configuration where the eccentricities of the two planets are locked at stationary points in the secular dynamics diagram. In this case, they show that the circularization time can be as high as $\\sim 8\\,$Gyr despite a standard $Q_\\mathrm{p}$ value.  To date the model by \\citet{bat09} is the only one able to explain the high present day eccentricity of \\glb\\ with a standard $Q_\\mathrm{p}$ value. This model could appear as not generic, as it requires a specific initial configuration, but \\citet{bat11} showed that in the framework of Hamiltonian planetary dynamics with additional dissipative forces (which is the case here), these points tend to behave as attractors. Hence various initial configurations can lead to reach such a point with a further dynamical evolution like described by \\citet{bat09}. As is shown below, in the model presented here, we exactly encounter such a configuration where many different routes can lead to such a stationary point (the transition between Phase 1 and Phase 2, see Sect. 4).  The purpose of this paper is to present an alternate model based on Kozai mechanism, assuming again a distant perturber. This model can be viewed as complementary to the \\citet{bat09} study. Kozai mechanism is a major dynamical effect in non-coplanar systems that can trigger eccentricity modulations up to very high values. We make a review on this effect, without and with tidal effects taken into account and describe our approach in Sects.~2 and 3 respectively. In Sect.~4, we present an application to the case of \\gl. We first present calculations starting from the present day orbit of \\glb, with the result that even Kozai mechanism cannot overcome the damping effect of tidal friction. Then we present other simulations where \\glb\\ is initially put much further away from the star than today. We show that Kozai mechanism pumps \\glb's eccentric regularly up to high values where tides become active at periastron. The result is a decay of the orbit that drives it to its present day location but on a time scale that can be considerably longer than the standard circularization time. Hence even after several Gyrs, \\glb's eccentricity can still be significant. In Sect.~5, we describe a parametric study of this scenario and derive in which conditions it is compatible with the present day situation of \\glb. Our conclusions are presented in Sect.~6. ",
        "conclusions": "The Kozai migration scenario is a generic process that causes inward migration of planets in non coplanar planetary systems. This evolution is characterized by two phases, a first one with perturbed Kozai cycles, and a second one with a more drastic decay of the semi-major axis. In any case, the characteristic time for this evolution can be considerably longer than the standard tidal circularization time, providing this way a solution to the issue of the high present eccentricity of \\glb. If we combine our study and that of \\citet{mon09}, various suitable initial configurations can be found. They all require a high initial mutual inclination ($\\ga 75\\degr$), an initial orbit for \\glb\\ several times wider than today, perturber masses ranging between less than $0.1\\,\\mjup$ and $50\\,\\mjup$, with orbital periods ranging between a few years up to several hundreds. This model is also generic enough to be invoked to explain some other puzzling cases of high eccentricity close-in exoplanets. But for a given set of initial conditions, the orbital period and the mass of the perturber must assume a fairly well constrained relationship, otherwise the suspected process is either too rapid or nonexistent. This could then explain why system like \\gl\\ are not numerous.  Note that the constraint derived for $i_0$ is valid for small initial eccentricity $e_0$. The real condition is $|h|\\la 0.3$, which can be achieved either with low $e_0$ and high $i_0$, but also with low $i_0$ and high $e_0$. For small $i_0$, $|h|\\la 0.3$ implies $e_0\\ga 0.95$. Although this is actually just a matter of fixing the starting point on the Kozai cycles, this has a different meaning for what concerns the formation of such a planetary system. Planetary systems are usually thought to form in circumstellar disk with initial coplanar configurations \\citep[see e.g. theoretical studies by][]{pie08,cri09}. The only way to imagine an initially inclined configuration is to assume that both planets formed independently and were assembled further to build the \\gl\\ system. This could happen for instance if \\gl\\ was initially member of a multiple stellar system. Such systems are indeed rarely coplanar. Alternatively, the system could have formed coplanar but the inner planet could have been driven to high eccentricity. There are many ways to achieve this. This can be done by planet-planet scattering events or chaotic diffusion. \\citet{lask96} indeed showed (with application to Mercury, see also \\citet{wu11}) that in any planetary system, the inner planet may be subject to drastic eccentricity increases due to secular chaos. Similarly, recent work on the so-called Nice model for the formation of the Solar System \\citep{tsig05,lev11} have demonstrated the role of planet-planet scattering among the giant planets. Mean-motion resonances can also drive inner bodies to high eccentricities. This mechanism was invoked to explain the suspected Falling Evaporated Bodies (FEB) scenario in the $\\beta\\:$Pictoris system \\citep{bm96,bm00}.  In conclusion, an initially nearly coplanar but eccentric configuration of the system appears more realistic than a highly inclined one. Note that this does not change anything to the further evolution described in this paper, as Kozai cycles cause the system to rapidly oscillate between these two extreme configurations. Even if the system is initially nearly coplanar, it evolves to highly inclined within half a Kozai cycle.  The relevant issue is actually the definition of the starting point.  According to \\citet{bat11}, in all systems potentially subject to Kozai resonance, Kozai cycles are first frozen as long as self gravity and strong planet-planet interactions induce a large enough apsidal precession. But at some point when the system relaxes, Kozai cycles can start, which constitutes the ``starting point'' in our scenario.  The mean-motion resonance model that applies to the FEB scenario can also provide a valuable starting point. The only requirement is that both planets are initially in a mean-motion resonance like 4:1 ou 3:1, with a small initial mutual inclination (a few degrees) and a small eccentricity ($\\ga 0.05$) for the outer, more massive planet. If this is fulfilled, the inner body is subject to an eccentricity increase virtually up to $\\sim 1$. It was shown by \\citet{bp3d} that whenever it reaches the high eccentricity state (the FEB state in the case of $\\beta\\:$Pictoris), it undergoes high amplitude inclination oscillations that can be viewed as Kozai cycles inside the mean-motion resonance. This is a kind of resonant version of our model that could constitute a full coherent scenario driving the system from a coplanar and circular configuration to strong Kozai resonance.  The reality of this scenario will nevertheless be difficult to test observationally. Even if the two planets are initially in mean-motion resonance, as soon the semi-major axis of \\glb\\ starts to evolve thanks to tidal dissipation, they leave the resonance. The present day orbital configuration keeps no memory of the initial resonance. Assuming that we are today in the middle of the second phase with a still significant eccentricity and an already reduced semi-major axis (by a typical factor 10), we are now unable to distinguish between initially resonant and non-resonant configurations. We also note from Fig.~\\ref{1a_long} that the present day mutual inclination between the two planets should be in any case $\\sim 20\\degr$, which makes it difficult also to distinguish observationally between coplanar \\citep{bat09} and Kozai models (this work). However, as explained above, the exact characterization of the apsidal behavior of the planets can lead to constraints on their mutual inclination. Basically, if $\\varpi-\\varpi'$ is close to 0 or $\\pi$, this would constitute a strong clue in favor of the coplanar model of \\citet{bat09}. Otherwize, this would indicate an inclined configuration.  Before that, a crucial issue is to continue the radial velocity and photometric follow up of \\gl. If we confirm the presence of a more distant companion, and if we are able to significantly constrain its orbit, then we will be able to test the proposed scenarii into more details. This would not only resolve the eccentricity paradox of \\glb, but it would also provide clues towards the system's past evolutionary history."
    },
    "1208/1208.2142_arXiv.txt": {
        "abstract": "{A collection of the best solar and laboratory spectra in the soft X-rays is used here to perform a preliminary  benchmark in this wavelength region, by comparing observed vs. predicted wavelengths and calibrated solar irradiances. The benchmark focuses on the \\ion{Fe}{ix} --  \\ion{Fe}{xiv} ions, for which we have recently calculated the relevant atomic data, however a few other ions have also been benchmarked. The iron ions are dominating the soft X-rays, however a large fraction of the strongest soft X-ray lines due to $n=4 \\to n=3$ transitions were previously unidentified. The strongest transitions are all identified here, in particular the decays from the core-excited levels (3s 3p$^l$ 4s, $l=$ 5,4,3,2,1 for \\ion{Fe}{x}, \\ion{Fe}{xi}, \\ion{Fe}{xii}, \\ion{Fe}{xiii}, and  \\ion{Fe}{xiv} respectively) which are the strongest soft X-ray transitions from these ions. Many new identifications are proposed, some only tentatively. Good agreement in terms of solar irradiances between the soft-Xray and EUV ($n=3 \\to n=3$) transitions is found, confirming the reliability of the new large-scale calculations. Some of the new atomic data and identifications are particularly important for the  Solar Dynamic  Observatory (SDO)  Atmospheric Imaging Assembly (AIA) 94~\\AA\\ band. ",
        "introduction": "The soft X-ray (50--170~\\AA) spectrum is rich in $n=4 \\to n=3$ transitions from  highly ionised iron ions, from \\ion{Fe}{viii} to \\ion{Fe}{xvi} (see, e.g.  \\citealt{fawcett_etal:72}, \\citealt{manson:72}, and \\citealt{behring_etal:72}). Various current missions are routinely observing the  soft X-rays. For example, Chandra with the LETG, and the  Solar Dynamic  Observatory (SDO) with a suite of instruments. The SDO  Extreme ultraviolet Variability Experiment (EVE) \\citep{woods_etal:12} has  been providing soft  X-ray irradiances long-ward of 60~\\AA, while the Atmospheric Imaging Assembly (AIA, see \\citealt{lemen_etal:12}) has been observing, for the first time routinely,  the solar corona in two broad-bands centred in the soft X-rays, around  94 and 131~\\AA.   Very little atomic data were available in the soft X-rays and the majority of the spectral lines still await firm identification. Within the APAP network (www.apap-network.org), we are carrying out a long-term project for calculating accurate atomic data for the soft X-rays. We started with the  \\ion{Fe}{viii}--\\ion{Fe}{xiv} iron ions. The atomic data for \\ion{Fe}{viii} and \\ion{Fe}{ix} have recently been discussed in \\cite{odwyer_etal:11_fe_9}, where new DW calculations for these two ions were presented. The main problems  related to calculating accurate atomic data  for the $n=4$ levels are discussed in \\cite{delzanna_etal:12_fe_10}, where new large-scale R-matrix atomic calculations for \\ion{Fe}{x} have been presented. A similar work on \\ion{Fe}{xi}, \\ion{Fe}{xii}, and \\ion{Fe}{xiii} has been presented in \\cite{delzanna_storey:12_fe_11, delzanna_etal:12_fe_12,delzanna_storey:12_fe_13}. New  atomic data for \\ion{Fe}{xiv} and \\ion{Fe}{xvi} have also recently been calculated with the R-matrix  method by \\cite{liang_etal:10_fe_14} and \\cite{liang_etal:09_na-like}.   It is therefore now possible to provide the first benchmark study for the soft X-rays for these iron ions, based on accurate atomic calculations. Previously,  \\cite{lepson_etal:02} provided some tentative identifications for \\ion{Fe}{vii} -- \\ion{Fe}{x} based on EBIT laboratory measurements and unpublished  distorted wave (DW)  calculations. \\cite{liang_zhao:10} discussed \\ion{Fe}{ix} -- \\ion{Fe}{xvi} emission lines using DW  calculations obtained with the Flexible Atomic Code (FAC) and Chandra LETG observations of Procyon. However, various problems with this work have been found. First,  almost all of their identifications were either previously known or are at odds with the present results. Second, large discrepancies between observed and predicted line fluxes were present. Third, the Procyon spectra were poor in terms of signal and spectral resolution, when compared to the solar spectra used in the present benchmark.   Recently, \\cite{testa_etal:12} also used Chandra LETG observations of Procyon to benchmark CHIANTI v.6 \\citep{dere_etal:97,dere_etal:09_chianti_v6} data, however no atomic data for the \\ion{Fe}{x} -- \\ion{Fe}{xiv} were available, with the exception of old (and incorrect)  DW scattering calculations  for  \\ion{Fe}{x}.    This paper is one in a series (see \\citealt{delzanna_etal:04_fe_10}, hereafter Paper~I) that aims to provide an assessment of  atomic data needed for the analysis of astrophysical spectra by benchmarking them against all available experimental data. The approach is observation-based, i.e. focuses  on the  brightest spectral lines that are observed in astrophysical  spectra. The paper is organised as follows. In Sect.~2, we give a brief review of previous observations we used for the benchmark. In Sect.~3 we present our results and in Sect.~4 we reach our conclusions. ",
        "conclusions": "This paper is the first benchmark for the soft X-ray lines. It is a summary of almost two years of work on the calculations and identifications of the  soft X-ray lines due to $n=4 \\to n=3$ transitions of the main iron ions. Large-scale  R-matrix and distorted wave  scattering calculations turned out to be both needed, to account for resonance enhancements in the excitation rates for the $n=4$ levels, and for cascading from higher levels.   The identification work proved very difficult, due to the lack of high-resolution well-calibrated spectra, the fact that the soft X-rays are notoriously packed with a large number of transitions from a range of ions, and that laboratory spectra and solar spectra are very much different.   The strongest iron transitions are all finally identified here. Very good agreement between the soft-Xray ($n=4 \\to n=3$) and EUV ($n=3 \\to n=3$) irradiances of the strongest lines is found for the first time, confirming the reliability of the new calculations.  In several cases, various discrepancies in the previous identifications have been found, and many tentative (new or revised) identifications have been proposed. Better experimental data and more atomic calculations on a range of other ions will be needed to confirm them. Some  calculations for other ions that produce strong lines in the soft X-rays  are already in progress.  With regard to the SDO AIA 94~\\AA\\ band, good progress has been made, with a new important identification of a strong  \\ion{Fe}{xiv} line at 93.61~\\AA, and the new calculations for  \\ion{Fe}{x}, \\ion{Fe}{ix} and \\ion{Fe}{viii}. At least one residual transition still need to be identified though.  The new  large amount of APAP atomic data  will be made available through the CHIANTI database, however this will require a new format and a new way to handle them. Work is in progress in this direction."
    },
    "1208/1208.1377_arXiv.txt": {
        "abstract": "Active region fan loops in AR 11076 were studied, in search of oscillations, using high cadence spectroscopic observations from \\textit{Extreme-ultraviolet Imaging Spectrometer} (EIS) on board {\\it Hinode} combined with imaging sequences from the \\textit{Atmospheric Imaging Assembly} (AIA) on board SDO. Spectra from EIS were analyzed in two spectral windows, \\FeXII\\ 195.12~\\AA\\ and \\FeXIII\\ 202.04~\\AA\\ along with the images from AIA in 171~\\AA\\ and 193~\\AA\\ channels. We find short ($<$3~min) and long ($\\approx$9~min) periods at two different locations. Shorter periods show oscillations in all the three line parameters and the longer ones only in intensity and Doppler shift but not in line width. Line profiles at both these locations do not show any visible blue-shifted component and can be fitted well with a single Gaussian function along with a polynomial background. Results using co-spatial and co-temporal data from AIA/SDO do not show any significant peak corresponding to shorter periods, but longer periods are clearly observed in both 171~\\AA\\ and 193~\\AA\\ channels. Space-time analysis in these fan loops using images from AIA/SDO show alternate slanted ridges of positive slope, indicative of outward propagating disturbances. The apparent propagation speeds were estimated to be 83.5 $\\pm$ 1.8~\\kms\\ and 100.5 $\\pm$ 4.2~\\kms, respectively, in the 171~\\AA\\ and 193~\\AA\\ channels. Observed short period oscillations are suggested to be caused by the simultaneous presence of more than one MHD mode whereas the long periods are suggested as signatures of slow magneto-acoustic waves. In case of shorter periods, the amplitude of oscillation is found to be higher in EIS lines with relatively higher temperature of formation. Longer periods, when observed from AIA, show a decrease of amplitude in hotter AIA channels which might indicate damping due to thermal conduction owing to their acoustic nature. ",
        "introduction": "Detection of waves and oscillations in the outer solar atmosphere helps us not only to understand their role in coronal heating but also to remotely diagnose the properties of the corona - coronal seismology (\\opencite{1970PASJ...22..341U}; \\opencite{1984ApJ...279..857R}). Active region oscillations had been reported by several authors, both from imaging (\\opencite{1999SoPh..190..249N}; \\opencite{2000A&A...355L..23D}; \\opencite{2001A&A...370..591R}; \\opencite{2003A&A...404L...1K}; \\opencite{2006A&A...448..763M}) and spectroscopic (\\opencite{1999A&A...347..355I}; \\opencite{2001A&A...368.1095O}, \\opencite{2002A&A...387..642O}, \\opencite{2009A&A...494..355O}; \\opencite{2008ApJ...681L..41M}, \\opencite{2010ApJ...713..573M}; \\opencite{2009ApJ...696.1448W}, \\citeyear{2009A&A...503L..25W}) observations. \\inlinecite{1999A&A...347..355I} and \\inlinecite{2001A&A...368.1095O} find oscillations with periods as short as 1~min using observations from CDS/SOHO. \\inlinecite{2002A&A...387..642O} observed oscillations over a sunspot region at different heights in the atmosphere with frequencies ranging from 5.4~mHz (185~s) to 8.9~mHz (112~s) using both CDS/SOHO and TRACE. \\inlinecite{2009A&A...494..355O} reported active region oscillations over a range of frequencies from 2~mHz (500~s) to 154~mHz (6.5~s) using the \\textit{Extreme-ultraviolet Imaging Spectrometer} (EIS) on board {\\it Hinode}. They also reported that the frequencies higher than 8~mHz are observed preferentially at the boundaries of bright plage regions. All of them interpreted these oscillations as either slow magneto-acoustic waves or fast mode waves. \\paragraph*{} Extended loop regions in the active region, were commonly found to host propagating disturbances, which enhanced the interest in studying these regions. Active region fan loops were widely studied using imaging observations, mainly from TRACE ({\\it e.g.} \\opencite{1999SoPh..187..261S}; \\opencite{2003A&A...404L...1K}; \\opencite{2000A&A...355L..23D}, \\opencite{2002SoPh..209...61D}, \\citeyear{2002A&A...387L..13D}) and EIT/SOHO \\cite{1999SoPh..186..207B}. Average values of the propagation speed, amplitude relative to the background and periodicities of these disturbances are 99.7$\\pm$3.9~\\kms, 3.7$\\pm$0.2\\% and 284$\\pm$10.4~s, respectively (see the review by \\opencite{2009SSRv..149...65D}). Simultaneous observations from TRACE and CDS/SOHO by \\inlinecite{2003A&A...404L..37M}, also reveal intensity oscillations of 5~min periodicity. \\inlinecite{2009A&A...503L..25W} reported 12~min and 25~min periods in these loops, both in intensity and velocity, using data from EIS/{\\it Hinode}. All these observed propagating disturbances are interpreted in terms of propagating slow magneto-acoustic waves which led to various applications in coronal seismology. There are also a few reports on outflows at the edges of active region which can contribute significantly to the slow solar wind (\\opencite{2007Sci...318.1585S}; \\opencite{2008ApJ...676L.147H}; \\opencite{2009ApJ...706L..80M}). Recently \\inlinecite{2010ApJ...722.1013D} pointed out that the oscillations in intensity and velocity, in some cases, are accompanied by in-phase oscillations in line width of the same period when the line profiles are fitted with a single Gaussian. They show that the presence of faint quasi-periodic upflows driven from below leading to an additional blue-shifted component in the line profile can cause oscillations in all the three line parameters. Hence, it is difficult to differentiate between waves and flows using only imaging data. \\inlinecite{2011ApJ...727L..37T} also reported propagating disturbances supporting the quasi-periodic upflow scenario using the simultaneous observations from XRT and EIS. These results highlight the need for simultaneous imaging and spectroscopic observations, to characterize the propagating disturbances properly. More recently, \\inlinecite{2011ApJ...737L..43N}, using data from EIS/{\\it Hinode}, reported the observation of propagating slow mode waves along the continuous outflow. They see an increasing correlation between intensity and velocity disturbances as we move away from the base of the outflow region. In this article, we report observations of oscillations with short ($<$ 3~min) and long periods ($\\approx$9~min) in an active region fan loop system, using simultaneous observations from the EIS/{\\it Hinode} and the \\textit{Atmospheric Imaging Assembly} (AIA) on board SDO. We also discuss the possible interpretation of these oscillations. ",
        "conclusions": "We searched for oscillations in active region fan loops using simultaneous high cadence spectroscopic (EIS/{\\it Hinode}) and imaging (AIA/SDO) data. We find two locations showing oscillations of short ($<$1~min to 3~min) and long periods ($\\approx$9~min). Short periods are observed with oscillations in all the three line parameters, intensity, Doppler shift, and line width which is possibly due to coupling of different MHD modes. Amplitudes of these oscillations, in all the line parameters, show significantly higher values in higher temperature line. Investigation at the same location from AIA does not show these periods with significant peaks. The amplitudes of these periods from AIA are very low and it is difficult to observe reliable oscillations with periods less than 2~min from AIA. This might be due to the increased background emission in the line of sight from the wider passbands of AIA. Longer periods identified over the fan loops show oscillations only in intensity and Doppler shift but not in line width. Space-time analysis from AIA over a loop region covering this pixel location, indicates oscillations with similar periods propagating with different speeds in different temperature channels. Recently, \\inlinecite{2012SoPh..279..427K} also find that velocities of the propagating disturbances in structures located at sunspot regions are temperature dependent. The apparent propagation speeds are 83.5$\\pm$1.8~\\kms\\ and 100.5$\\pm$4.2~\\kms\\ in 171~\\AA\\ and 193~\\AA\\ channels, respectively. The ratio of these speeds, is close to the theoretical acoustic speed ratio in this region. This might indicate that these oscillations are due to slow magneto-acoustic waves. These waves also show lower amplitudes in the hotter channel of AIA consistent with damping due to thermal conduction. Line profiles at both these locations do not show any visible blue-shifted emission component. The asymmetry in the line profiles due to any possible periodic upflows can be masked if the field lines are highly inclined to the line of sight \\cite{2012ApJ...748..106T}. But the observation of significant Doppler shift oscillations at these locations might indicate that this is not the case here. We require detailed spectroscopy to completely resolve the ambiguity and the forthcoming IRIS mission can be very useful to address this.  \\begin{acks} We would like to thank the referees for their useful comments. Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSC (Norway). The AIA data used here is the courtesy of SDO (NASA) and AIA consortium. \\end{acks}"
    },
    "1208/1208.3658_arXiv.txt": {
        "abstract": "Cosmological measurements  require the calculation of nontrivial quantities over  large datasets. The next generation of survey telescopes  will yield measurements of billions of galaxies. The scale of these datasets, and the nature of the calculations involved, make cosmological calculations ideal models for implementation on graphics processing units (GPUs). We consider two cosmological calculations, the two-point angular correlation function and the aperture mass statistic, and aim to improve the calculation time by constructing code for calculating them on the GPU. Using CUDA, we implement the two algorithms on the GPU and compare the calculation speeds to comparable code run on the CPU. We obtain a code speed-up of between 10 - 180x faster, compared to performing the same calculation on the CPU. The code has been made publicly available. GPUs are a useful tool for cosmological calculations, even for datasets the size of current surveys, allowing calculations to be made one or two orders of magnitude faster. ",
        "introduction": "\\label{sec:intro}  The use of graphics processing units (GPUs) in scientific computing has been steadily growing in fields as diverse as bioinformatics, QCD lattice calculations and seismology (see, for example,~\\cite{bio, qcd, seis}). In astronomy, GPUs have proven useful in many different  computationally intensive problems such as N-body simulations (\\cite{2012JCoPh.231.2825B, nitadori}) and radio astronomy measurements (\\cite{clark}). GPU techniques have succeeded in reducing compute times for these difficult calculations by up to a factor of 100, and in this work we show that a similar reduction can be achieved for the calculation of cosmological quantities.  The next generation of large-scale astronomical surveys (such as the Dark Energy Survey~\\footnote{http://www.darkenergysurvey.org}, PanSTARRS~\\footnote{http://pan-starrs.ifa.hawaii.edu/}, and  the Large Synoptic Survey Telescope~\\footnote{www.lsst.org}) will produce enormous amounts of data, with measurements of billions of stars and galaxies. The problems associated with processing such a large volume of image data and how to structure and access a database containing tens of petabytes of information has been studied and discussed at length in the literature (e.g. ~\\cite{way, berriman, brunner}). In this paper, we address the challenges an astronomer will face attempting to analyze this information, once it has been obtained from a central database. With such a large volume of data the statistical uncertainties on many cosmological quantities will be reduced by orders of magnitude compared to present limits, but in order to take advantage of this data new computation methods must be developed. Methods used today to calculate cosmological quantities tend to be of complexity $O(n^{\\ 2})$ (where $n$ is the number of data points), which is not computationally feasible with the billions of measurements expected from future surveys, even with the expected improvements in computer hardware.  Many cosmological calculations require independent calculations of the same quantity for all data points, which makes them ideal candidates for parallelization (where each calculation is farmed out to a different processing device). In the past this has often been handled by using many CPUs in a computing cluster environment, but building these clusters or buying time on these systems can be expensive for researchers. GPU computing has brought a significant amount of this computational power to the desktop and is therefore more affordable for individual analysts. As GPU computing develops and becomes more widely used by the broader astronomy community, we foresee performing calculations in the near future that are currently computing-limited. In this paper, we describe the GPU implementation of two of these cosmological calculations, the two-point angular correlation function and aperture mass statistic, using the CUDA programming language (\\cite{cuda}).  Large-scale structure in the universe is a valuable probe of the composition and evolution of matter in the universe, and  can be used to constrain models of cosmology. Galaxies are good tracers of the total matter in the universe; although we can only directly detect the luminous matter, galaxies form around concentrations of dark matter. We can therefore characterize large scale structure of matter in the universe using the clustering of galaxies on different length scales, which is measured using the angular correlation function. The two-point angular correlation function (or matter power spectrum in Fourier space) is based on galaxy number counts, and measures the excess or depletion of pairs of galaxies as a function of separation, compared to a random distribution. At small scales  ($\\approx 100 h^{-1}Mpc$) we can measure the imprint of the baryon acoustic oscillations, which gives information on the phases of acoustic waves at recombination. At larger scales the matter power spectrum has not been affected by radiation or baryons, so it is a rare probe of primordial fluctuations and inflation. See~\\cite{bassett} for a full review of the subject, \\cite{peebles} for a description of large-scale structure in the universe and~\\cite{cole, eisenstein} for the first measurements of baryon acoustic oscillations. Calculating a correlation function over billions of galaxies, and at separations ranging from arcseconds to degrees, requires significant computational power, which scales with the square of the number of galaxies. As such it is an excellent candidate for implementation on the GPU. Higher order correlation functions (such as the three-point correlation function) can also be used to probe non-Gaussian features in the galaxy distribution. These calculations are even more computationally expensive than the two-point correlation function, and would potentially benefit enormously from parallelization on the GPU.   As we were preparing this paper, we became aware of work by~\\cite{ponce} presenting a method and code for the calculation of the two-point correlation function on the GPU using CUDA. We take a different approach to the implementation of the algorithm. We have also become aware of earlier work by~\\cite{Roeh}, which also takes another different approach to the implementation, and also considers MPI implementation and the use of multiple GPUs in parallel.  Dark matter cannot yet be detected directly, but the effect of its gravitational field can be measured indirectly using gravitational lensing. When light travels through the universe, its path is deflected by the gravitational potential of the matter it passes. The distortion of the observed shapes of distant galaxies as their light passes through matter in the universe encodes information about large-scale structure and the growth of matter in the universe. However, since galaxies have intrinsic shapes, it is only through statistical analysis of large numbers of galaxies that we can average out their intrinsic shapes and orientations and extract the cosmological information. There are many ways to interpret this information (see~\\cite{bs} for a review), but for the purposes of this paper we concentrate on the shear peak statistic. Background galaxies tend to have the major axis tangential to foreground mass density contours, so concentrations of matter along the line of sight can be detected from the coherent distortions in the ensemble of their shapes. Over-densities will appear as peaks in a map of aperture mass. Counting the number of peaks as a function of peak significance allows us to distinguish between different cosmologies (\\cite{marian09, kratochvil09, dh09}). We calculate the aperture mass at a point by making a weighted sum over the tangential components of the ellipticity of the surrounding galaxies. Typically this sum contains tens or hundreds of thousands of galaxies, and must be performed on a dense grid of points over the sky in order to accurately reconstruct the projected mass field. The sum for each point is independent, and can be performed in parallel making it another good candidate for implementation on the GPU. Recent work by ~\\cite{leonard} has used wavelet transformations to speed the calculation of the aperture mass, but as far as we are aware this is the first implementation on the GPU.  This paper is structured as follows. In Section~\\ref{sec:ang} we introduce the two-point angular correlation function and describe its implementation on the GPU. We then compare its performance to the same calculation on the CPU using data taken from dark matter simulations. In Section~\\ref{sec:sps} we describe the aperture mass calculation and how it is implemented on the GPU and compare its performance to the CPU implementation. In Section~\\ref{sec:summary} we summarize our findings. Access to the code is given in \\ref{app:coderep}. ",
        "conclusions": "\\label{sec:summary} We have implemented code to perform calculations of the two-point angular correlation function and the aperture mass statistic on the GPU. We have demonstrated that this implementation can reduce compute times for these calculations by factors of 100$\\times$-300$\\times$, depending on the amount of data to be processed. The code for making this calculation is publicly available from Github (see \\ref{app:coderep} for details). With a straightforward division of the dataset into sub-sets, our code can also be used on clusters of independent GPUs. Faster compute speeds mean that a full MC-based calculation of the errors for the angular correlation function can reasonably be performed, without approximations or assumptions that are required to make the calculation reasonable for CPU codes (e.g. kd-tree calculation (\\cite{Jarvis:2003wq})). We intend to evaluate this in future work.  The increasing size of astronomical datasets will require a new approach to data analysis. We expect that the use of the GPU in everyday cosmological calculations will become more common in the next few years, especially since faster compute times allows experimentation in techniques used to make the calculation and rapid comparison of the results. We expect this application to be extended to other computationally challenging calculations, such as the three-point and higher order angular correlation functions, and the shear correlation functions.     \\appendix"
    },
    "1208/1208.4849_arXiv.txt": {
        "abstract": "The existence of an extended hot gaseous corona surrounding clusters, groups and massive galaxies is well established by observational evidence and predicted by current theories of galaxy formation. When a small galaxy collides with a larger one, their coronae are the first to interact, producing disturbances that remove gas from the smaller system and settle it into the corona of the larger one. For a Milky-Way-size galaxy merging into a low-mass group, ram pressure stripping and the Kelvin-Helmholtz instability are the most relevant of these disturbances. We argue that the turbulence generated by the latter mixes the material of both coronae in the wake of the orbiting satellite creating a ``warm phase'' mixture with a cooling time a factor of several shorter than that of the ambient intragroup gas. We reach this conclusion using analytic estimates, as well as adiabatic and dissipative high resolution numerical simulations of a spherical corona subject to the ablation process of a constant velocity wind with uniform density and temperature. Although this is a preliminary analysis, our results are promising and we speculate that the mixture could potentially trigger in situ star formation and/or be accreted into the central galaxy as a cold gas flow resulting in a new mode of star formation in galaxy groups and clusters. ",
        "introduction": "The classical theory of galaxy formation establishes that gas follows dark matter in its non-linear evolution as it forms virialized dark matter haloes. In this process, the gas is heated by shocks and adiabatic compression, acquiring the virial temperature of the halo it has fallen into \\citep{White_Rees_1978}. Afterwards, the gas can lose energy effectively through radiative cooling if the cooling time is much smaller than the free-fall time ($t_{\\rm cool}<t_{\\rm ff}$), producing a contraction of the gas towards the centre: a cooling flow that fuels star formation. The critical mass for effective cooling depends on the density, temperature and metallicity of the gas. Assuming constant gas density and temperature, one can show that the gas in haloes with primordial gas masses $\\gtrsim10^{11}$M$_{\\rm \\odot}$ cannot cool effectively \\citep[e.g. see Fig. 8.6 of][]{Simon_book_2010}; this is why galaxy groups and clusters are expected to have a corona of hot gas at the present time.  X-ray observations of relaxed clusters and groups show indeed the presence of an extended gas corona within the intra-cluster and intra-group medium (hereafter called ICM in both cases), which is approximately in hydrostatic equilibrium within the gravitational potential of the cluster/group. For the most massive systems, the density profile of this diffuse gas approximately traces that of the predicted dark matter distribution \\citep[e.g.][]{Vik_2006}. Since  the dominant cooling channel in these systems is bremsstrahlung radiation, then $t_{\\rm cool}\\propto T^{1/2}\\rho^{-1}$ (for primordial composition) which implies that the cooling time scales are typically shorter in the central regions, and hence, in the absence of a balancing heat source, a cooling flow would develop in the centre. Close to the virial radius of the system however, the typical cooling times are larger than the Hubble time.  A hot extended corona is also predicted to exist within the haloes of massive galaxies, including those at the scale of the Milky-Way \\citep[MW;][]{White_Frenk_1991}. In the absence of additional energy sources, and under idealized circumstances, the majority of this gas would be expected to cool and collapse into the halo centre over a Hubble time; in contrast, roughly $20\\%$ of the expected baryons in these haloes are observed to be in a cold phase (stars and cold gas, see e.g. \\citealt{Fukugita_2004,Fukugita_2006}, and references therein). This overcooling problem can be avoided by invoking feedback processes that prevent the gas from cooling catastrophically. Recent cosmological simulations that include radiative cooling, star formation and feedback predict that MW-size haloes should have a gas corona surviving today in quasi-hydrostatic equilibrium \\citep[e.g.][]{Crain_2010,voort_2011}. It has also been argued that even without feedback, the combined effect of buoyancy and thermal conduction suppresses the thermal instability, and thus the formation of cold clouds, in stratified coronae in MW-size galaxies \\citep{Binney_2009, Nipoti_2010}.  The presence of extended coronae has been confirmed for elliptical galaxies through X-ray observations \\citep[see review by][]{Mathews_2003} but the evidence for MW-like galaxies is not conclusive. Recent observations have detected diffuse X-ray emission around disc-dominated galaxies, but is approximately two orders of magnitude lower than the expectations based on the \\citet{White_Frenk_1991} model \\citep[e.g.][]{Li_2007}; supernovae heating is advocated as the likely cause of this discrepancy \\citep[e.g.][]{Crain_2010,Anderson_2011,Dai_2011}. We note that although there is evidence for the presence of gaseous coronae around galaxies, their cosmological origin as a shock-heated remnant is not settled. There are other possible origins such as enriched gas produced by galactic stellar feedback \\citep[e.g.][]{Ciotti_1997,Mathews_1998}.  It is thus relevant to consider in detail the impact of the gaseous coronae on the process of galaxy formation, particularly during galaxy mergers. In these events, the gas in the merging galaxies is subject to different processes that remove gas from the satellite: i) ram pressure stripping removes the most loosely bound material \\citep{GG_1972} creating a wake that trails behind the satellite (recent observations show evidence of such wakes in galaxies falling into clusters, e.g. \\citealt{Sun_2010}, but the stripped gas in the observed cases is the cold ISM of the galaxy, not its corona); ii) thermal conduction causes evaporation of the colder gas in the satellite \\citep{Cowie_1977}; iii) if the flow around the boundary interface is laminar then viscosity will cause a drag that will remove gas \\citep{Nulsen_1982}; iv) fluid instabilities such as the Kelvin-Helmholtz (KH) instability caused by the velocity shear between the layer interface also lead to gas removal.  All these processes have been studied in the past and for the case of galaxies merging into clusters, particular attention has been given to ram pressure stripping using analytic and numerical methods \\citep{Mori_2000,Heinz_2003,McCarthy_2008}. Most of these studies focus on the impact of the removal of gas in the star formation efficiency of the satellite: by completely stripping the gas corona, a fuel reservoir is removed, resulting in a quenching of star formation relative to a case where the galaxy evolves in isolation (this process is usually called ``strangulation''). Recent analyses show that a substantial fraction of the corona can actually survive if the satellite is massive enough, and through subsequent radiative cooling, can contribute to star formation \\citep{McCarthy_2008,Moster_2011}.  The fate of the stripped gas has not been analysed in detail however. As mentioned before, the gas from the satellite removed by ram pressure will form an elongated wake trailing the satellite. Due to the gravity of the satellite, a similar wake will be formed with gas from the ICM that gets gravitationally focused \\citep{Stevens_1999,Sake_2000}. Thus, in general we would expect a wake composed of material from both the ICM and the corona of the satellite, implying that both gaseous components mix during the removal processes. This mixing is particularly efficient during the turbulence caused by KH instabilities. If the gas removal rate is large and the mixing process is efficient then the mixture will have a temperature that might be significantly lower than the temperature of the ICM. This would imply that the cooling time in the wake is lower than that in the ambient medium, making the mixture a potential fuel source for star formation. This idea has been proposed recently by \\citet{Marinacci_2010} in the context of a mode of gas accretion triggered by the interaction of H$_{\\rm I}$ clouds expelled from galactic discs in a galactic fountain, and the corona of the galaxy \\citep[see also][]{HP_2009}. In this paper we argue that the same mechanism is at play when the coronae of the merging galaxies interact.  Most of the hydrodynamical simulations of galaxy formation carried out to date have used the smoothed particle hydrodynamics (SPH) technique \\citep[for a review see][]{Springel_2010}. This is a formulation that has excellent conservation properties in terms of mass, energy, momentum and entropy and it is fully Galilean invariant. However, SPH considerably suppresses the formation of turbulence triggered by instabilities and thus it is not able to account for gas mixing \\citep{Agertz_2007,Bauer_2011}. Adaptive Mesh Refinement (AMR) codes based on an Eulerian formulation are the other method currently used for simulations of galaxy formation. Although they suffer their own shortcomings (like the lack of Galilean invariance), perhaps the most important difference with SPH methods is precisely that of the treatment of turbulence and mixing \\citep{Mitchell_2009}, which are not suppressed in AMR.  In this paper we take a first step in investigating the effect of mixing in the interaction of coronae in galaxy mergers, focusing on the formation of the trailing wake. For this purpose we start by revisiting analytically the most important gas removal processes and estimate the properties of the resultant gas mixture (Section \\ref{theo}). We then verify these estimates using 3D adiabatic (non-radiative), as well as dissipative simulations of the motion of a spherical corona through a constant density ICM using the AMR code FLASH (Section \\ref{sims}). Our results are presented in Section \\ref{results} and finally a discussion and our conclusions are given in Sections \\ref{sec_discussion} and \\ref{conclusions}. Although our simulations lack several important physical processes that could impact the results, our objective here is to make an initial assessment of whether or not the removed coronal gas could continue to be a source of star formation. ",
        "conclusions": "The presence of an extended corona of hot gas within clusters, groups and massive galaxies is a generic prediction of the current theory of galaxy formation. Since such coronae are a large gas reservoir that can potentially cool down and form stars, it is of great relevance to study their evolution and in particular their interaction during galaxy mergers which are ubiquitous in the CDM paradigm.  Most studies of galaxy mergers have concentrated on the interaction of the cold components (stars and gaseous discs), relegating the hot coronae to a secondary role essentially governed by ram pressure stripping. In the case of MW-size galaxies merging into groups, the high temperature and low density of the ambient ICM at the outskirts of the groups are commonly assumed to imply that the stripped gas stays irrelevant for subsequent cooling and star formation. This assumption, however, ignores the fact that Kelvin-Helmholtz instabilities produced by the velocity shear also remove coronal gas and efficiently mix it with the ICM through turbulent eddies. Since the former is colder than the latter, the mixture could have a shorter cooling time than the ambient intragroup gas. These instabilities are typically suppressed in the majority of galaxy formation simulations that have been done to date since most of them were carried out using the SPH formulation \\citep{Agertz_2007}.  In this paper, we use analytic methods and hydrodynamical simulations (with and without radiative cooling) to study the gas removal processes that act during the interaction a of a spherical MW-size corona moving through a homogeneous ICM wind representative of the coronal gas at the virial radius of a low-mass group. The main conclusions of our work are the following: \\begin{itemize} \\item Our simulations show that a relevant fraction of the removed coronal gas mixes with the ICM in trailing wakes creating a ``warm phase'' (i.e. with at least 50\\% less entropy than the ambient ICM) with a cooling time which is almost an order of magnitude lower than the cooling time of the ambient intragroup gas. This reduction is a result of a lower temperature in the gas mixture driven by the colder coronal gas from the satellite, as anticipated, but also due to an enhanced density caused by compression. \\item We find that the gas mass in this warm phase is between 40-76\\% of the mass stripped from the satellite and deposited in the wakes ($M_{\\rm warm}\\sim1.4-7.6\\times10^9$M$_\\odot$ for the parameters we have used in our simulations). This gas is mostly coming from the satellite, but a non-negligible fraction, 20-28\\%, comes from the ICM. \\item The analytic model described in Section \\ref{theo} is able to roughly estimate the amount of removed mass and the time scale of removal due to the two main acting mechanisms: ram pressure stripping and KH instabilities. Moreover, once properly calibrated, it can be used to estimate the most likely temperature and cooling time of the gas mixture after the gas settles in the turbulent wakes. We believe that this model can be incorporated into semi-analytic models of galaxy formation to account for these processes, which are typically ignored in current implementations. \\end{itemize}  We speculate that the warm phase created through the mechanism proposed in this work has the potential to become a mode of star formation in galaxy groups (through in situ star formation and/or through a cold stream flowing to the central galaxy of the host). In fact, the possibility of stars forming in turbulent wakes from the condensation of gas stripped from cold discs seems to be a likely hypothesis in some observed cases \\citep[e.g.][]{Hester_2010}. Although in such cases the mechanism at work is the same we have studied here, star formation is far more likely in those since the gas being removed is much colder than the ambient gas. We are thus cautious in speculating about the fate of the removed coronal gas. In particular, it is not clear whether the gas mixture can survive long enough without evaporating through conduction with the ICM \\citep[e.g.][]{Nipoti_2004}. This would depend on the specific details of conduction in the surfaces of the turbulent wake, which are difficult to estimate. It is therefore crucial to address these questions using more detailed simulations. We believe that the results of this work motivate further analyses to fully quantify the effect of the mixing of the coronal gas from satellites with the intragroup gas."
    },
    "1208/1208.4880_arXiv.txt": {
        "abstract": "We report on the results of a direct $N$-body simulation of a star cluster that started with $N = 200\\,000$, comprising $195\\,000$ single stars and $5\\,000$ primordial binaries. The code used for the simulation includes stellar evolution, binary evolution, an external tidal field and the effects of two-body relaxation. The model cluster is evolved to $12\\,$Gyr, losing more than 80\\% of its stars in the process. It reaches the end of the main core-collapse phase at $10.5\\,$Gyr and experiences core oscillations from that point onwards -- direct numerical confirmation of this phenomenon. However, we find that after a further $1\\,$Gyr the core oscillations are halted by the ejection of a massive binary comprised of two black holes from the core, producing a core that shows no signature of the prior core-collapse. We also show that the results of previous studies with $N$ ranging from $500$ to $100\\,000$ scale well to this new model with larger $N$. In particular, the timescale to core-collapse (in units of the relaxation timescale), mass segregation, velocity dispersion, and the energies of the binary population all show similar behaviour at different $N$. ",
        "introduction": "\\label{s:intro}  The increasing ability of the direct $N$-body method to provide reliable models of the dynamical evolution of star clusters has closely mirrored increases in computing power (Heggie 2011). The community has progressed from small-$N$ models performed on workstations (e.g. von Hoerner 1963; McMillan, Hut \\& Makino 1990; Giersz \\& Heggie 1997) to models of old open clusters and into the $N \\sim 100\\,000$ regime (Baumgardt \\& Makino 2003) by making use of special-purpose GRAPE hardware (Makino 2002). Software advances over a similar timeframe have produced sophisticated codes (Aarseth 1999; Portegies Zwart et al. 2001) that increase the realism of the models by incorporating stellar and binary evolution, binary formation, three-body effects and external potentials. As a result, $N$-body models have been used in numerous ways to understand the evolution of globular clusters (GCs: Vesperini \\& Heggie 1997; Baumgardt \\& Makino 2003; Zonoozi et al. 2011), even though the best models still only touch the lower end of the GC mass-function (see Aarseth 2003 and Heggie \\& Hut 2003 for a more detailed review of previous work). At the other end of the spectrum, Monte Carlo (MC) models have proven effective at producing dynamical models of $N > 10^6\\,$particles (Giersz \\& Heggie 2011). These models have shown that clusters previously defined as non-core-collapse can actually be in a fluctuating post-core-collapse phase (Heggie \\& Giersz 2008). In practice the two methods are complimentary with MC informing the more laborious $N$-body approach (such as refining initial conditions) and $N$-body calibrating aspects of MC.  In this paper we present an $N$-body simulation of star cluster evolution that begins with $N = 200\\,000$ stars and binaries. This extends the $N$ parameter space covered by direct $N$-body models and performs two important functions. Firstly it provides a new calibration point for the MC method -- this statistical method is increasingly valid for increasing $N$ so calibrations at higher $N$ are more reliable. It also allows us to further develop our theoretical understanding of star cluster evolution and investigate how well inferences drawn from models of smaller $N$ scale to larger values. The latter is the focus of this current paper. A good example of the small-$N$ models that we wish to compare with is the comprehensive study of star cluster evolution presented by Giersz \\& Heggie (1997) using models that included a mass function, stellar evolution and the tidal field of a point-mass galaxy, albeit starting with $500$ stars instead of $200\\,000$. More recent examples for comparison include Baumgardt \\& Makino (2003) and K\\\"{u}pper et al. (2008). We were also motivated to produce a model that exhibited core-collapse close to a Hubble time without dissolving by that time. What we find when interpreting this model is that much of the behaviour reported previously for smaller $N$-body models stands up well in comparison but that the actions of a binary comprised of two black holes (BHs) provides a late twist to the evolution of the cluster core.  In Section 2 we describe the setup of the model. This is followed by a presentation of the results in Sections 3 to 7 focussing on general evolution (cluster mass and structure), the impact of the BH-BH binary, mass segregation, velocity distributions and binaries (binary fraction and binding energies). Throughout these sections the results are discussed and compared to previous work where applicable. Then in Section 8 we specifically look at how the evolution timescale of the new model compares to findings presented in the past. ",
        "conclusions": "\\label{s:discus}  We have presented an $N$-body model that started with $200\\,000$ stars and binaries, evolves to the moment of core-collapse at $10.5\\,$Gyr and has $\\sim 30\\,000$ stars remaining at $12\\,$Gyr. We have used our direct $N$-body model to confirm the post-core-collapse fluctuations described in the Monte Carlo model of Heggie \\& Giersz (2008) and the hybrid $N$-body/MC approach of Heggie \\& Giersz (2009). We have also shown that these fluctuations can be halted by the ejection of a dominant BH-binary from the core. This produces a core that shows no sign that it has previously evolved through core-collapse.  We have looked at how the results of previous works compare to a model of larger $N$ and find good agreement provided that appropriate scalings are used (such as the core-radius to half-mass radius ratio at core-collapse). In terms of raw values some variations exist: the core radius at core-collapse reaches deeper in to the mass distribution for larger $N$, for example. The behaviour of quantities such as average stellar mass and velocity dispersion have been documented and the general behaviour matches expectations from earlier models. Looking at time scales such as the time to core-collapse and the dissolution time we also find agreement with scaling relations previously reported in the literature, however this is dependent on which value of the half-mass relaxation timescale is used. In particular, the scaling of dissolution time with $t_{\\rm rh}^{3/4}$ reported by Baumgardt (2001) could be reproduced provided that the average $t_{\\rm rh}$ was used and not the initial $t_{\\rm rh}$.   The $N = 200\\,000$ simulation reported in this work took the best part of a year on a GRAPE-6 board to complete. It continues the gradual increase of $N$ used in realistic $N$-body models from the $N = 500$ model of Giersz \\& Heggie (1997) to the $N = 131\\,072$ model of Baumgardt \\& Makino (2003). However, with only $\\sim 20\\,000 \\, M_\\odot$ remaining in our model at an age of $11.5\\,$Gyr we are still only touching the lower end of the globular cluster mass-function. This is after considerable effort. In particular, we are still some way from the goal of a full million-body model of a globular cluster (Heggie \\& Hut 2003). How can we push forward to reach that goal? The shift towards graphics processing units (GPUs) as the central computing engine for $N$-body codes, combined with sophisticated software development, offers hope (Nitadori \\& Aarseth 2012). Simulations of $100\\,000$ stars can be performed comfortably on a single-GPU (Hurley \\& Mackey 2010; Zonoozi et al. 2011) and the introduction of multiple-GPU support will likely make simulations of the type presented here commonplace in the near future. Further hardware advances and a revisiting of efforts to parallelize direct $N$-body codes (Spurzem 1999) will also aid the push towards greater $N$.  In a follow-up paper we will conduct a full investigation of the stellar and binary populations of our model. It is also our intention to make model snapshots -- saved at frequent intervals across the lifetime of the simulation -- available for others to {\\it observe} and analyse. These can be obtained by contacting the authors. At the beginning of this paper we indicated that an important function of a large-$N$ model would be to aid in the calibration of the Monte Carlo technique. This is currently underway (Giersz et al. 2012) and will include a direct comparison of $N$-body and MC models starting from the same initial conditions."
    },
    "1208/1208.4622_arXiv.txt": {
        "abstract": "SCUBA-2 is the largest submillimetre wide-field bolometric camera ever built. This 43 square arc-minute field-of-view instrument operates at two wavelengths (850 and 450 microns) and has been installed on the James Clerk Maxwell Telescope on Mauna Kea, Hawaii. SCUBA-2 has been successfully commissioned and operational for general science since October 2011. This paper presents an overview of the on-sky performance of the instrument during and since commissioning in mid-2011. The on-sky noise characteristics and NEPs of the 450$\\mu$m and 850$\\mu$m arrays, with average yields of approximately 3400 bolometers at each wavelength, will be shown. The observing modes of the instrument and the on-sky calibration techniques are described. The culmination of these efforts has resulted in a scientifically powerful mapping camera with sensitivities that allow a square degree of sky to be mapped to 10 mJy/beam rms at 850$\\mu$m in 2 hours and 60 mJy/beam rms at 450$\\mu$m  in 5 hours in the best weather.\\\\ ",
        "introduction": "\\label{sec:intro}  SCUBA-2 is a dual-wavelength camera containing 5120 pixels in each of two focal planes. Each focal plane consists of 4 separate rectangular sub-arrays each with 1280 bolometers butted together to give the full field of approximately 8 $\\times$ 8 arcmin.  The two wavebands, which target the atmospheric windows centred at 450$\\mu$m and 850$\\mu$m (as was its predecessor, SCUBA~\\cite{holland1999}),  observe the sky simultaneously by means of a dichroic beamsplitter. SCUBA-2 was delivered from the UK Astronomy Technology Centre to the Joint Astronomy Centre in Hawaii in April 2008 with one engineering sub-array at each waveband. The focal-planes were fully-populated with science-grade arrays in summer 2010 and the first astronomical data with the full complement of arrays was taken in early-2011.  Submillimetre observations are more critical than ever before in a broad range of astronomical areas of interest. SCUBA-2 is presently the largest submillimetre camera in the world, and the fastest wide-area ground-based submillimetre imager. An ambitious set of legacy surveys is currently well underway and these are designed to maximise the impact of SCUBA-2's wide-field, fast-mapping capabilities and to complement the current and future set of submillimetre facilities, such as Herschel and ALMA. The surveys include large-scale surveys of  molecular clouds and star-formation regions (Gould-Belt Survey and Galactic Plane Survey), an extensive survey of nearby galaxies (NGLS), targeted study of debris discs around nearby stars (SONS), a deep cosmology survey (S2CLS) and a broad survey of the outer Galaxy (SaSSy).  This paper will outline the performance of SCUBA-2 as it undertakes its first semester of full science observations, following the successful commissioning of the instrument in October 2011. The instrument characterization results will be presented, including its noise performance on the sky, responsivity and bolometer yields. The observing strategies will be discussed, and the final instrument sensitivities will be shown. Finally, a selection of the first science observations will be presented. A detailed article on the instrument performance is given in Ref.~\\citenum{holland2012}.\\\\ ",
        "conclusions": "SCUBA-2 has been commissioned and operational for science observations since October 2011. The phases of on-sky commissioning were discussed, including the flat-fielding methodology and absolute-power calibration. The noise performance of the instrument was presented for the first science semester, beginning in February 2012. The NEPs show good stability over both short and long time scales, though some unknown noise features are observed and are under investigation. The atmospheric calibration technique was outlined and the waveband-dependent extinction correction terms were shown. The flux calibration has been successfully characterised, with absolute deviation in the FCF performance of less than 10\\% at both wavelengths. Description of the observing modes and data-reduction were given. Finally, the instrument sensitivities were calculated, showing that in the best weather, a 1-square-degree observation reaches 10 mJy sensitivity in approximately 2hrs at 850$\\mu$m, while at 450$\\mu$m, the same size map can reach 60 mJy rms in 5 hours.\\\\"
    },
    "1208/1208.0417_arXiv.txt": {
        "abstract": "We have calculated the physical conditions throughout the merger Seyfert galaxy NGC 6240 by modelling the observed optical and infrared line ratios. We have found that the optical spectra are emitted by clouds  photoionised by the power-law radiation flux from  the AGN (or AGNs), and heated mainly by the shock accompanying the propagation of the clouds outwards. The infrared line ratios are emitted from  clouds ejected from a starburst which photoionises the gas by the black-body radiation flux corresponding to a stellar colour temperature  of about 5 10$^4$ K. Both the flux from the AGN and the ionization parameters are low. The most characteristic physical parameters  are the relatively high shock velocities ($\\geq$500 \\kms) and low preshock densities ($\\sim$ 40-60 \\cm3) of the gas. The   C/H, N/H, O/H  relative abundances  are higher than solar by a factor $\\leq$ 1.5. We suggest that those relative abundances indicate   trapping of  H into H$_2$ molecules rather than  high metallicities. Adopting an initial grain radius of 1 \\mum, the dust temperatures calculated in the clouds reached by the power-law  radiation and by the black-body radiation are 81 K and 68 K, respectively. ",
        "introduction": "Colliding galaxies are  at first identified  by the morphological appearance of the encounter, depending on the  merging evolution age, but they  are hardly   recognizable  through the physical properties of the new galaxy. Eventually, the spectra show the dominant role of dynamical processes such as shock waves and starbursts created   by collision, yet the active galactic nuclei (AGNs) have a dominant role.  Evidences of galaxy collision  in the Seyfert galaxy NGC 6240  appeared on prints of the Palomar Sky Survey (Fosbury \\& Wall 1979, hereafter FW79), namely, a complex high luminosity structure with  substantial plumes extending over 130 kpc,  large dust lanes and chaotic appearance of the nuclear region. The synchrotron characteristic of the radio emission, the coincidence of a central feature seen in the optical (\\Ha + [NII] emission) with radio emission, the strength of the low level ionization lines reveal the presence of  shock waves as a result of collision. FW79 in fact suggested   that the NGC 6240 system is the result of an encounter between galaxies. They  dismissed the hypothesis that the  collision  could trigger star formation,  comparing the observed line ratios with  those emitted from gas ionized by   the black body radiation flux from young stars.  On the other hand, Rieke et al. (1985) defined  NGC 6240 as a site of exceedingly powerful bursts of star formation, involving nearly 10$^{10}$ \\msol of newly formed stars. This was suggested not only by the luminosity and extent of the stellar component near 2 \\mum, but particularly on the basis of the large population of red supergiants required to account for the depth of the stellar CO absorption.  NGC 6240 was  definitively  classified among mergers by Fried \\& Schulz (1983) who detected a \"pronounced\" double nucleus on an S1 image tube plate in the I band and on r and I exposures taken with a CCD camera. The two nuclei are oriented at a position angle of 10$^o$ and separated by 1.8\" (640-700 pc, projected on the sky) with diameters  of 2.5\" and 1.5\".  Fried \\& Schulz claimed that both nuclei appear to be non-stellar,  rejecting the hypothesis of an accidental projection of a foreground star. The extended regions observed by  them are dominated by shock heating on a scale of $\\sim$ 6 kpc. The binary active nucleus in NGC 6240 was confirmed  by Komossa et al (2003). Using CHANDRA, they discovered two hard X-ray nuclei coincident with the optical-IR ones. Each nucleus host an AGN detected in hard X-rays and at 5 GHz (Gallimore \\& Beswick 2004).  A starburst in NGC 6240 is the source of the K-band luminosity, dominated by supergiants (Tecza et al. 2000). The starburst was triggered quite recently and has  a duration time  less than 5 million yrs. It  fills  the central kpc, the region encompassing the two nuclei (Engel et al 2010). Heckman  et al. (1990) showed  evidence of a  starburst driven super-wind on a large scale ($>$ kpc).   At a distance of 97 Mpc (z=0.0245,  H$_o$=75 \\kms Mpc$^{-1}$), NGC 6240 has an infrared luminosity L$_{IR}$$\\sim$ 10$^{11.8}$ \\Lsol (Engel et al  2010). It was included in the list of ultra-luminous infrared galaxies (ULIRGs) by Rigopoulou et al. (1999). The dynamical field is characterised by the extraordinary large stellar velocity dispersions ($\\sim$ 350 \\kms). The optical spectrum of NGC 6240 shows many emission lines with line ratios  characteristic of a collisional dominated regime rather than   a power-law radiation dominated flux  generally  prevailing in the narrow line region (NLR) of AGN (FW79).  The mid-infrared (IR) spectrum taken with the Infrared Spectrograph on the {\\it Spitzer Space Telescope} by Armus et al (2006) shows  the flux of the most significant strong fine-structure lines up to [SiII] 34.8 \\mum, as well as rotational H$_2$ lines and PAH emission features.   Armus et al  inquired  about the  buried AGN contribution to the bolometric luminosity. The mid-IR spectrum was included in the survey of the AGN NLR  in the IR by Dasyra et al (2011).  Comparing mid-infrared emission-line properties from high-resolution Spitzer spectra of a hard X-ray (14-195 keV) selected sample of nearby (z $<$ 0.05) AGNs detected by the Burst Alert Telescope (BAT) aboard Swift, Weaver et al. (2010) found that the luminosity distribution  of the [O IV] 25.89 \\mum, [Ne II] 12.81 \\mum, [Ne III] 15.56 \\mum, and [Ne V] 14.32, 24.32 \\mum lines and  hard X-ray continuum are similar for  Seyfert 1 and Seyfert 2 populations. They also found  that the emission lines primarily arise in gas ionized by the AGNs. On the other hand, Luhman et al. (1998) report measurements of the [C II] 157.74 \\mum fine-structure line in ULIRGs with the Long Wavelength Spectrometer on the Infrared Space Observatory. They claim that the observation of [CII]/ FIR continuum ratio is consistent with starburst nuclei.  In this paper we investigate the source of radiation and the physical conditions of the emitting gas in NGC 6240 by a detailed modelling of  the optical spectrum presented by FW79 and of the mid-IR spectrum   reported by Armus et al.  The clumpy and irregular  morphology of NGC 6240 and the relatively large FWHM of the spectral lines indicate that, besides the photoionization flux from the AGNs, shocks  are heating and ionizing the NLR gas. The leading role of shocks was already noticed by FW79 on the basis of the strong neutral and low level ionization lines. Indeed, different regions throughout  merging galaxies show  physical conditions changing on small scales (e.g. for NGC 7212 and NGC 3393). The spectrum observed from NGC 6240 is an average of the spectra emitted from the different regions. Trying to disentangle  various components throughout a unique spectrum leads to  approximated results. However, the rather peculiar line ratios observed  by FW79 suggest that the task is worthy. We will refer to the observed spectra with the maximum precision as we have done for  other merger galaxies (e.g. NGC 3393, Contini 2012,  NGC 7212, Contini et al 2012, MKN 298, Radovich et al. 2005), i.e. trying to find some records of  collision  analysing the spectra. A first hint can be obtained by  modelling  the spectral line ratios and the spectral energy distribution (SED) of the continuum. Our aim, in particular, is to  distinguish  the contribution  of the AGN from that of the  starbursts in a  strongly dominated shock wave regime.  We will adopt composite models which account for the flux from the active centre (AC) and/or  for  radiation from the stars, consistently coupled with shock wave hydrodynamics. We have found in previous investigations (e.g. Rodr\\'{i}guez-Ardila et al. 2005, Contini et al. 2004a, Viegas et al. 1999, and references therein) that,  not only  shocks  are important to explain high ionization level lines (e.g. [FeVII], [FeX]) coronal lines (e.g. [SiVIII], [SiIX])  as well as low ionization and neutral lines (e.g. [OI] and [NI]), but also they determine the intensity and the frequency of the SED peaks in the X-ray  and IR domains.  We have run a grid of models in order to select the most appropriated one reproducing the NGC 6240 observed line ratios. Building the grid by many models we could explore which ranges of the physical parameters could characterise  the  NLR of NGC 6240.  The  calculations of the spectra are presented in Sect. 2. The optical spectrum presented by FW79  and the fine-structure line ratios in the IR observed by Armus et al. are investigated in Sect. 3. The calculated continuum SED is compared with the observational data in Sect. 4. Discussion and concluding remarks follow in Sect. 5. ",
        "conclusions": "The Seyfert galaxy NGC 6240 is the result of  collision and  merging. Merging  objects  generally include 1) a double active nucleus, 2) stars and starbursts in the central region and in the edges, even beyond them, and 3) shocked fragmented clouds. Therefore, the  models adopted for the calculation of the spectra   account for the power-law radiation flux from the AGNs,  black-body radiation flux from the starburst and shock waves.  We have calculated the physical conditions throughout the  galaxy by modelling the optical, infrared line ratios and  the continuum SED. The best fitting models were selected from a large grid.  \\subsection{The physical picture}  Modelling suggests that  the clouds photoionised by both the AGN  flux and by the starburst  move outwards, i.e. they are not gravitationally falling towards the black holes. The starburst region embraces the zone between the AGNs and beyond it and a starburst super-wind carries the clouds outwards (Heckman 1990).  We have found that   the flux from  the AGN (or AGNs) $F$ is  similar to the lowest  ones found in LINERs (e.g. Contini 1997) and LLAGNs (e.g. Contini 2004a). It leads to  optical line ratios   best fitting the data. The pl radiation flux photoionises the clouds and heats the gas at maximum temperatures of 3-4 10$^4$ K. The gas is  mainly heated by the shock to maximum temperatures of 3.75 10$^6$ K and 1.2 10$^7$ K accompanying  the clouds outwards with shock velocities \\Vs $\\sim$ 500-900 \\kms, respectively. A high \\Vs combined with a low \\n0  ($\\sim$ 40 \\kms) leads to the high [OII]/[OIII] observed in NGC 6240.  A distance r $\\sim$ 370 pc from the   AC is calculated for the  gas reached by the AGN radiation and emitting the observed optical line ratios, combining the \\Hb absolute flux calculated at the nebula by   model m2$_{pl}$ (Table 1) with \\Hb observed at Earth : H${\\beta}_{abs}$ r$^2$= H${\\beta}_{obs}$ d$^2$, where d is the distance from Earth and H${\\beta}_{obs}$=1.7 10$^{-14}$ \\erg (FW79). The distance r  is a lower limit. A higher distance should result accounting for the filling factor, \\ff=10$^{-3}$ - 10$^{-4}$ (Heckman et al 1990). The results obtained for r by modelling the continuum in Sect. 4 lead to \\ff= 0.074.  The infrared spectrum reveals the starburst. We have found that the  IR line ratios are emitted from gas photoionised by a starburst corresponding to a stellar colour temperature \\Ts  $\\sim$ 5 10$^4$ K. The gas is included in   geometrically thick  clouds  ($D$= 3 pc). The ionization parameter  is low ($U$=0.002), indicating that the starburst region is extended and   the photon flux is  diluted on its way to the emitting gas.  In Fig. 5 we plot  [NeV]/[NeII] versus [NeIII]/[NeII] for both the Type I and Type II AGN samples of Dasyra et al (2011). We chosed only Ne lines in order to avoid the relative abundance problem. Fig. 5 shows  that the two samples almost coincide, suggesting that the AGN is not the main source of ionization. Indeed the starbursts and the shocked gas throughout the NLR dominate, as  also appears by comparing NGC 6240  with the AGN samples  in Fig. 5.  \\begin{figure*} \\includegraphics[width=14.cm]{fg5.eps} \\caption{[NeV]/[NeII] versus [NeIII]/[NeII] for the sample of Type 1 and Type 2 AGN presented by Dasyra et al. (2011) : blue squares and black circles, respectively. The red circle represents NGC 6240; the red asterisk indicates the starburst model result. Red  + and red X represent the AGN models calculated by \\Vs=500 \\kms and 900 \\kms, respectively. } \\label{fig3} \\end{figure*}  The  clouds emitting the [NeII] 12.8 lines are at a distance r $\\geq$ 4.3 kpc from the starburst, which is calculated by combining the [NeII] absolute flux calculated by model m1$_{sb}$ (Table 1) with [NeII] observed  at Earth  (Dasyra et al 2011) : [NeII]$_{abs}$ r$^2$= [NeII]$_{obs}$ d$^2$, where d is the distance of NGC 6240 to Earth. A distance of 4.3 kpc was found comparing the calculated to the observed continua, leading to \\ff $\\leq$ 1.  The models presented in Table 1 are constrained enough to allow to calculate the location of the [CII] emitting region. Combining the flux  observed at Earth with that calculated at the nebula, we obtain r= 9 kpc  for the AGN (by model m2$_{pl}$) and 11 kpc for the starburst (by model m2$_{sb}$), adopting a filling factor of 1.   We found in agreement with FW97 that the preshock densities are relatively low ($<$ 100 \\cm3) compared to those calculated in the NLR of Seyfert galaxies ($>$ 100 \\cm3), but  higher than  those of the ISM ($>$ 1 \\cm3). The preshock densities are  relatively low  in  both the clouds ionised by  the starburst and by the AGN. The characteristic low density in the NLR of NGC 6240  suggests that  the gas from the ISM was inflated inside the colliding galaxies  and mixed up, diluting their former densities. Alternatively, both the colliding galaxies were characterised by  low density gas similar to  that  found near the centre of the Milky Way.  The shock velocities are relatively high  (\\Vs=500-900 \\kms) and similar for  most of the clouds which are carried by the wind,  confirming that the high shock velocities arised from an  extended collision event. They are slightly higher  than the observed FWHM of the lines ($\\sim$ 400 \\kms), indicating that shocks   are stronger than the bulk velocity of the clouds in some positions.  The  similarity of some of the physical conditions predicted by the  AGN and  starburst models (e.g. low densities, high velocities and  relative abundances slightly higher than solar) is peculiar, indicating that  the clouds are  most probably  remnants of shells ejected  from the starburst. The matter swept up  by the blast wave accelerates outward and fragments via Rayleigh-Taylor instabilities. Relatively  small geometrical thickness of the clouds ($D$ = 0.03 - 0.2 pc) are calculated in the neighbourhood of the AGNs. However, they are large enough  to include a large zone of neutral gas (Fig. 3). Here, the electron temperature and densities are kept relatively high by the secondary radiation, so the neutral emission lines ([NI] and [OI]) are high enough to fit the observed [NI]/\\Hb and [OI]/\\Hb line ratios. The clouds which better reproduce the infrared line spectrum are ionised by the starburst and have  geometrical thickness ($D$= 3 pc) higher by a factor $\\geq$ 10. Neutral line ratios are lower than in the clouds close to the AGNs,  due to a large  region of cold gas inside the cloud. Different geometrical thickness are in agreement with an underlying turbulence (Contini \\& Goldman 2011)  \\subsection{Relative abundances of the  heavy elements in the gaseous phase}  The influence of merger-induced inflows, enrichment and gas consumption,  galactic winds, etc. is  important in determining the nuclear metallicity. The central metallicity is primarily a competition between the inflow of low-metallicity gas and enrichment from star formation (Torrey et al 2012).  We have  obtained the best fit to the spectra by  C/H= 1. - 1.3 solar, N/H=1.5 solar,  O/H= 1.45 solar and Ar is twice solar. The relative abundances derived from  our modelling of the NGC 6240 emission spectra are not common in mergers at low z. Before considering  enrichment for instance from star formation, let us  consider  some  evidences about grains and molecules in a collision dominated regime such as that in NGC 6240.  Tacconi et al (1999) report the result of the CO J=2$\\rightarrow$1 line observations. Half of the CO flux is concentrated in a thick disc structure situated between the two radio/infrared nuclei, with a diameter $\\leq$ 500 pc. Strong CO bands were already found in NGC 6240 by Rieke et al (1985) indicating giant and supergiants, where dust is formed. Also H$_2$ v=1$\\rightarrow$0 emission is centered between the radio/infrared nuclei (Joseph et al 1984).  Mac Law \\& Glover (2011) and Glover \\& Mac Law, (2012) calculations on the formation of H$_2$ and CO in clouds, show  that the CO molecules are easily dissociated for low densities ($\\sim$ 30 \\cm3) while H$_2$ molecules are still  present,  partially removing H from the gaseous phase. Strong molecular hydrogen emission has been detected by Rieke et al (1985) and by Becklin et al (1984) in NGC 6240.  The 1-0S(1)H$_2$ emission line is one of the most powerful found in any galaxy to date (Joseph et al 1984), probably excited by shocks (Tecza et al 2000, etc.), while CO abundance around the nuclei is very likely to have been significantly reduced by X-ray irradiation from the AGN. It seems that  the H$_2$ emission is excited purely by thermal mechanisms   (Sugai et al 1997) and H$_2$ lines are emitted between the double nuclei. From its excitation mechanism and its peak position,   the emission is from a global shock caused by a galaxy-galaxy collision.  We suggest that   abundances of heavy elements  relative to H are high in NGC 6240 due to trapping of hydrogen into H$_2$ molecules, removing  it from the gaseous phase. High  C/H, N/H,   O/H and Ar/H are therefore an indirect record of  collision, They do not indicate  high metallicities, but they reveal the formation of H$_2$ molecules in a large collision event. We cannot predict the  true C/H, N/H and O/H relative abundances, perhaps lower than solar. Yet, the remnants of the  inflated gas can explain Heckman et al discovery of two very bright knots located about 8.5 kpc west and 2.5 kpc WSW of the nucleus. They have line widths less than 100 \\kms FWHM and \\Ha is about 15 times brighter than [SII] 6717+ and at least 25 times brighter than [NII] 6584. Heckman et al. claim that these knots are probably giant HII regions with \\Ha uncorrected luminosities of about 4  10$^{39}$ and 9  10$^{39}$ erg s$^{-1}$, respectively and that  [NII] and [SII] lines suggest very low metal abundances, less than 20\\% solar. We suggest that they represent the original inflowed matter, not processed by the super-wind.  On the other hand, elements which are easily trapped into grains and molecules are strongly depleted from the gaseous phase. Fe is generally included into small grains and therefore easily sputtered. The depletion of Fe indicates  that depleted matter was included at collision of the galaxies. S and Si can be included into molecules, but downstream of shocks at \\Vs=500-900 \\kms everything is evaporated. So they were also  originally depleted in the included matter, in agreement with Heckman et al."
    },
    "1208/1208.5004_arXiv.txt": {
        "abstract": "Using simulations of slowly rotating stratified turbulence, we show that the $\\alpha$ effect responsible for the generation of astrophysical magnetic fields is proportional to the logarithmic gradient of kinetic energy density rather than that of momentum, as was previously thought. This result is in agreement with a new analytic theory developed in this paper for large Reynolds numbers. Thus, the contribution of density stratification is less important than that of turbulent velocity. The $\\alpha$ effect and other turbulent transport coefficients are determined by means of the test-field method. In addition to forced turbulence, we also investigate supernova-driven turbulence and stellar convection. In some cases (intermediate rotation rate for forced turbulence, convection with intermediate temperature stratification, and supernova-driven turbulence) we find that the contribution of density stratification might be even less important than suggested by the analytic theory. ",
        "introduction": "Turbulent dynamos occur in many astrophysical situations. They tend to develop large-scale magnetic structures in space and time that are generally understood in terms of mean-field dynamo theory \\citep[e.g.][]{Mof78,Par79,KR80,ZRS83,RSS88,RH04,BS05}. Central to this theory is the $\\alpha$ effect, which denotes a contribution to the mean electromotive force that is given by a pseudo-scalar $\\alpha$ multiplying the mean magnetic field. Such a pseudo-scalar can be the result of rotation, $\\OO$, combined with stratification of density and/or turbulence intensity, $\\nab\\meanrho$ and/or $\\nab\\urms$, respectively. Here, $\\meanrho$ is the mean gas density and $\\urms=(\\overline{{\\bm u}^2})^{1/2}$ is the rms value of the turbulent velocity, ${\\bm u}$.  There have been a number of analytic studies quantifying the effects of rotating stratified turbulence on the mean electromotive force \\citep{KR80,K91,RK93,KPR94,RKR03,KR03}. In particular, it was found, using the quasi-linear approach (or second-order correlation approximation), that the diagonal components of the $\\alpha$ tensor for slow rotation rate (or small Coriolis numbers) are given by \\citep{SKR66,KR80} \\begin{eqnarray} \\alpha\\approx-\\ell_\\alpha^2\\OO\\cdot\\nab\\ln(\\meanrho \\, \\urms), \\label{I1} \\end{eqnarray} where $\\ell_\\alpha=\\tau_0\\urms$ is a relevant length scale and $\\tau_0$ is the characteristic turbulent time related to the turnover time.  In the solar convective zone the mean fluid density $\\meanrho$ changes by seven orders of magnitudes, while the turbulent kinetic energy ($\\approx \\meanrho \\urms^2/2$) changes by only three orders of magnitudes and, according to stellar mixing length theory \\citep{Vit53}, $\\meanrho \\, \\urms^3$ would be approximately constant in the solar convective zone. Here the fluid density $\\rho$ is the sum of the mean and fluctuating density. This issue has become timely, because there is a new numerical technique that allows the different proposals to be examined with sufficient accuracy. The so-called test-field method \\citep{Sch05,Sch07} allows one to determine all the relevant turbulent transport coefficients in the expression for the mean electromotive force without the restrictions of some of the analytic approaches such as the quasi-linear approach, the path-integral approach or the $\\tau$ approach. With the test-field method one solves sets of equations for the small-scale fields resulting from different prescribed mean fields --- the test fields. These equations resemble the usual induction equation, except that they contain an additional inhomogeneous term.  This method is quite powerful because it has been shown to be rather accurate and it gives not only the tensor coefficients of $\\alpha$ effect and turbulent diffusivity, but it also allows the scale-dependence to be determined, which means that these coefficients are actually integral kernels that allow the effects of neighboring points in space and time to be taken into account.  For details regarding scale separation, see \\cite{BRS08} and \\cite{HB09}.  We apply this method to numerical simulations of forced turbulence in a stably stratified layer in the presence of rotation and a prescribed vertical dependence of the turbulence intensity.  We also use test-field method in simulations of turbulent convection and supernova-driven turbulence of the interstellar medium (ISM).  The goal of the present paper is to determine the correct scaling of the $\\alpha$ effect with mean density and rms velocity for slow rotation, i.e., when the Coriolis number ${\\rm Co}\\equiv 2 \\Omega \\tau_0$ is much less than unity, and large Reynolds numbers. In addition to the parameter $\\ell_\\alpha$, we determine the exponent $\\sigma$ in the diagonal components of the $\\alpha$ tensor, \\begin{eqnarray} \\alpha = - \\ell_\\alpha^{2} \\OO \\cdot \\nab \\ln(\\meanrho^\\sigma\\urms). \\label{I2} \\end{eqnarray} Such an ansatz was also made by \\cite{RK93}, who found that in the high conductivity limit, $\\sigma=3/2$ for slow rotation and $\\sigma=1$ for rapid rotation. However, as we will show in this paper, both numerically (for forced turbulence and for turbulent convection with stronger temperature stratification and overshoot layer) as well as analytically, our results for slow rotation and large fluid and magnetic Reynolds numbers are consistent with $\\sigma=1/2$. In some simulations we also found $\\sigma=1/3$. ",
        "conclusions": "While the present investigations confirm the old result that the $\\alpha$ effect in mean-field dynamo theory emerges as the combined action of rotation and stratification of either density or of turbulent intensity, they also now point toward a revision of the standard formula for $\\alpha$. The old formula by \\cite{SKR66} predicted that the effect of stratification can be subsumed into a dependence on the gradient of $\\meanrho\\urms$. This formula was then generalized by \\cite{RK93} to a dependence on $\\meanrho^\\sigma\\urms$, where $\\sigma=3/2$ in the high conductivity limit for slow rotation, and $\\sigma=1$ for faster rotation. In contrast, our new results now clearly favor a value of $\\sigma$ below unity. The idealized case of artificially forced turbulence can most directly be compared to our analytic derivation, since it agrees in all the made assumptions. The obtained value of $\\sigma=1/2$ agrees very well with the theoretical expectation. A similar exponent is found for the case of turbulent convection with higher temperature stratification, but the results seem to depend sensitively on model parameters (see~\\Tab{summaryall}). Here more detailed studies will be required. Moreover, the result $\\sigma=1/2$ arises naturally from analytical considerations for large fluid and magnetic Reynolds numbers and slow rotation as the only tenable choice, but those considerations have not yet been performed for the cases of intermediate and rapid rotation.  Forced turbulence simulations show a trend toward smaller values of $\\sigma$ around 1/3 for faster rotation and also in cases of supernova-driven turbulence. Turbulent convection with overshoot also gives 1/3 in one case of moderate temperature stratification with overshoot, while simulations without overshoot point toward values somewhat larger values around 3/4. However, in none of the cases we have found that the $\\alpha$ effect diminishes to zero as a result of a trend toward constant convective flux for which $\\meanrho\\urms^3$ is approximately constant. In spite of the considerable scatter of the values of $\\sigma$ found from various simulations, it is worth emphasizing that in all cases $\\sigma$ is well below unity. On theoretical grounds, the value 1/2 is to be expected. Except for the forced turbulence simulations that also yield 1/2 for slow rotation, all other cases are too complex to expect agreement with our theory that ignores, for example, inhomogeneities of the density scale height and finite scale separation."
    },
    "1208/1208.5468_arXiv.txt": {
        "abstract": "\\noindent We present deep surface photometry of a volume--limited sample of $21$ UM emission line galaxies in broadband optical $UBVRI$ and near infra-red (NIR) $HKs$ filters. The sample comprises $19$ blue compact galaxies (BCGs) and two spirals. For some targets the exposure times are the deepest to date. For the BCG UM462 we observe a previously undetected second disk component beyond a surface brightness level of $\\mu_B=26$ mag arcsec${}^{-2}$. This is a true low surface brightness component with central surface brightness $\\mu_0=24.1$ mag arcsec${}^{-2}$ and scale length $h_r=1.5$ kpc. All BCGs are dwarfs, with $M_B\\ge-18$, and very compact, with an average scale length of $h_r\\sim1$ kpc. We separate the burst and host populations for each galaxy and compare them to stellar evolutionary models with and without nebular emission contribution. We also measure the $A_{180}$ asymmetry in all filters and detect a shift from optical to NIR in the average asymmetry of the sample. This shift seems to be correlated with the morphological class of the BCGs. Using the color-asymmetry relation, we identify five BCGs in the sample as mergers, which is confirmed by their morphological class. Though clearly separated from normal galaxies in the concentration--asymmetry parameter space, we find that it is not possible to distinguish luminous starbursting BCGs from the merely star forming low luminosity BCGs. ",
        "introduction": "\\noindent Blue compact galaxies (BCGs) are low metallicity gas--rich galaxies at low redshifts, currently undergoing intense star formation (SF). Their star formation rates (SFR) are on average too high to be indefinitely sustained by the available gas supply. Their spectra are reminiscent of \\HII regions, with strong emission lines superposed on a blue stellar continuum, which is why they are sometimes referred to as \\HII galaxies. Deep optical and near infra--red (NIR) observations have revealed the presence of an old stellar population in these galaxies, often referred to as the ``host'', in which the starbursting regions are embedded. The original criteria of what constitutes a BCG~\\citep{1981ApJ...247..823T} referred to compactness ($r_{25}\\sim1$ kpc in diameter) on photographic plates, blue colors, and low total luminosity ($M_B\\gtrsim-18$), however, with the discovery of an old and extended underlying host population in almost all BCGs~\\citep[e.g.][]{1996A&AS..120..207P,1997MNRAS.286..183T,2001ApJS..133..321C,2001ApJS..136..393C,2002A&A...390..891B,2003ApJ...593..312C,2003A&A...410..481N}, these criteria have been relaxed to be more inclusive. Thus, BCGs comprise a heterogeneous group of galaxies, with varied morphologies, star formation histories, and total luminosities, but they all have \\HII region emission line spectra, which is in practice their only unifying characteristic.\\\\ \\subsection*{Sample selection} \\noindent This paper is part of a series and should be read as such. In~\\citet[][hereafter \\M]{Paper1} we presented and analysed $UBVRIHKs$ broadband imaging for a sample of 24 BCGs. That sample was hand--picked to contain interesting and representative cases of BCGs and is biased towards relatively luminous (median $M_B\\sim-18$ mag) galaxies. The \\M~ sample is defined in terms of galaxy class -- all galaxies are BCGs -- but the heterogeneous and hand--picked nature of the sample make it difficult to translate the properties of such an inherently mixed bag of BCGs to global properties of the galaxies in the local Universe. An inherent problem is that the BCG classification is somewhat ad hoc and based on criteria mainly relating to their appearance on photographic plates rather than their star forming properties, and most samples have ill determined completenesses. In an attempt to study a spatially well defined sample of BCGs complete in terms of luminosity we turned to emission line surveys. In magnitude limited surveys a galaxy's inclusion in the survey depends entirely on its apparent brightness, which introduces a bias against low luminosity systems despite the fact that those are the most common ones. Understandably, one would like to study the most common type of galaxy in the Universe which makes emission line surveys, with their small/no luminosity bias, a favorable place to look for a representative and abundant sample of such systems.~\\citet{1989ApJS...70..447S} compiled a large sample of emission line galaxies (ELGs) from Lists $IV$ and $V$ of the University of Michigan (UM) objective-prism survey. The primary selection criteria for this survey are based on the strength and contrast of the [OIII] $\\lambda5007$ emission line and it therefore contains a larger fraction of low luminosity dwarfs compared to magnitude limited surveys~\\citep{1989ApJS...70..479S}. BCGs, being a subgroup of emission line galaxies, make up about two thirds of the UM survey~\\citep{1989ApJS...70..479S}. Our approach in this paper is to take a volume of space and study all emission line galaxies in it. We use~\\citet{1989ApJS...70..447S} to select a volume limited sample defined by $11\\le RA\\le14$h and $v\\le2100$ km s${}^{-1}$. This velocity cut--off ensures that we have good completeness at the faint end~\\citep[][completeness $\\gtrsim95\\%$ for $v<2500$km s${}^{-1}$]{1989ApJ...347..152S}. Inside of this volume are $21$ UM ELGs, of which $19$ are BCG--like and two are giant spiral galaxies. Thus selected, this sample is representative of the star forming galaxy population in the local Universe. It consists predominantly of compact low--luminosity dwarfs of various (burst) metallicities -- from low ($Z\\sim0.004$) to close to solar ($Z\\sim0.02$), and with varying gas content. Throughout this paper we refer to the sample galaxies, with the exception of the two spirals, as BCGs.\\\\  \\noindent These galaxies and the targets from \\M~together constitute a sample of $46$ high and low luminosity BCGs. The observations presented here are a part of our ongoing effort to study representative numbers of such galaxies. Kinematic data exist for the majority and are about to be published (\\\"Ostlin et al. 2012 in prep., Marquart et al. 2012 in prep.).~\\citet{1997MNRAS.288...78T} find that such galaxies readily divide into two major morphological types (roughly into regular and irregular), which indicates that they may have different progenitors. The deep optical and NIR imaging data in this paper and in \\M~will allow us to study the difference in the faint old populations of these two groups and compare their structural parameters and photometric properties. Though we will frequently refer to the properties of the BCGs in \\M~throughout this paper, the bulk of the analysis juxtaposing low and high luminosity BCGs will appear in a dedicated future paper~\\citep{Paper3}. We have assumed $H_0=73~km^{-1}s^{-1}Mpc^{-1}$.\\\\  \\noindent The layout of this paper is as follows: \\S~\\ref{data} introduces the data and the calibration, and provides a log of the observations. \\S~\\ref{methods} briefly summarizes the derived profiles and the measured quantities. \\S~\\ref{individ} gives brief notes on the characteristics of the galaxies, as well as a detailed summary of how stellar evolutionary models (SEMs) compare with the observed colors for each galaxy. Where possible, an indication of the age and metallicity for the different populations is given. Observed trends in the integrated colors, asymmetries, total luminosities, and other galaxy properties are discussed in \\S~\\ref{discuss}. We summarize our conclusions in \\S~\\ref{conclude}. \\section[]{Observations}\\protect\\label{data} \\noindent The data consist of optical and NIR broadband imaging, obtained during the period $2003$--$2007$ with ALFOSC (at the Nordic Optical Telescope, NOT), MOSCA (NOT), and EMMI (at the European Southern Observatory New Technology Telescope, ESO NTT) in the optical, and with NOTCAM (NOT) and SOFI (ESO NTT) in the NIR. \\\\  \\noindent We have presented in detail the reduction pipelines and the calibration of the data in~\\citet{2010MNRAS.405.1203M} and \\M. We shall not repeat it here, except to give some brief notes on MOSCA reductions since \\M~did not contain any such data.\\\\  \\noindent \\textbf{MOSCA} is a multi-chip instrument (4 CCDs). Each CCD had its own illumination gradient, which was not aligned in concert with the rest. Since our pipeline fits and subtracts a sky from every reduced frame before stacking, it became necessary to adapt it to fit 4 separate skies and subtract those from the individual CCDs on each reduced frame, instead of fitting a single sky on the mosaiced (raw) frames. Dark current (DC) frames were available, however after subtracting the masterbias from the masterdark we found the remaining DC to be negligible for our longest exposure of 10 minutes, hence we did not use the DC frames in the reduction. The MOSCA bias level can occasionally fluctuate throughout the night on some of the CCDs, but again, after examination of the bias frames taken on five separate occasions throughout the night we found it to be very stable for all 4 CCDs. This is not universally the case with MOSCA data, so care must be taken to check the behavior of the dark and bias levels in the four chips for each individual night. The orientation of the four chips is slightly misaligned, which we have corrected for before making the final stacked images. \\\\  \\noindent We should further mention that during the reductions of this sample we again made extensive use of the \\textit{astrometry.net} software~\\citep{2010AJ....139.1782L} to add a world coordinate system (WCS) to the headers of most of the NOTCAM data. Possibly useful for the community tips, derived from our experience with this software, can be found in \\M.\\\\  \\noindent Tables~\\ref{exptable} and~\\ref{nedtable} summarize the individual exposure times for each filter and the observation log for these data. The heliocentric redshift and distance in Mpc, both taken from NED\\footnote{\\footnotesize NASA/IPAC Extragalactic Database, http://ned.ipac.caltech.edu/}, are also provided. The filter number at the respective observatory is given for convenience. The sample is volume--limited, with $11\\le RA\\le 14$h and $v\\le2100$ km s${}^{-1}$.\\\\ \\setcounter{table}{0} \\begin{table} \\begin{minipage}{70mm} \\caption{Total integration times for the sample. All times are given in minutes and converted to the framework of a 2.56 meter telescope where needed. The values are for observations in a single filter, e.g. only $SOFI~Ks$, and not $SOFI~Ks+NOTCAM~Ks$. }\\protect\\label{exptable} \\begin{tabular}{|lccccccc|} \\hline &U&B&V&R&I&H&Ks\\\\\\hline\\hline UM422&20&60&30&$30^\\dagger$&58&23&$125^\\ddagger$\\\\\\hline UM439&60&40&40&9&58&&249\\\\\\hline UM446&&40&40&9&116&21&51\\\\\\hline UM452&70&60&20&7&46&37&64\\\\\\hline UM456&40&40&40&7&38&30&32\\\\\\hline UM461&60&50&40&6&38&32&82\\\\\\hline UM462&30&40&40&6&38&32&48\\\\\\hline UM463&30&40&40&&38&30&73\\\\\\hline UM465&40&40&50&&&&121\\\\\\hline UM477&40&20&20&6&38&&62\\\\\\hline UM483&20&35&30&9&48&&249\\\\\\hline UM491&60&40&40&9&58&&249\\\\\\hline UM499&40&36&36&&19&&62\\\\\\hline UM500&60&60&30&$30^\\dagger$&38&61&123\\\\\\hline UM501&40&40&40&6&38&32&121\\\\\\hline UM504&30&40&40&9&9&31&121\\\\\\hline UM523&10&40&40&6&38&&62\\\\\\hline UM533&20&40&40&6&38&&80\\\\\\hline UM538&30&30&20&9&58&&145\\\\\\hline UM559&60&40&40&9&12&7&123\\\\\\hline \\hline \\end{tabular} \\medskip ~\\\\ $\\dagger$ -- ALFOSC, $\\ddagger$ -- SOFI \\end{minipage} \\end{table} \\subsection[]{Photometric calibration}\\protect\\label{photcal} \\noindent All data were calibrated in the Vega photometric system. We remind the reader that the calibration in the optical was carried out with Landolt standard stars, while in the NIR we used 2MASS to calibrate against the mean zero point of field stars found in each individual frame, which makes the NIR calibration less dependent on photometric conditions. In the optical we compared the photometry of stars in our calibrated frames with SDSS photometry in the same fields. Any offset larger than 0.05 mag detected between our photometry and the SDSS photometry was then applied to our frames. For both wavelength regimes we estimated the zero point uncertainty, $\\sigma_{zp}$, for each final frame as the average residual difference in magnitudes between SDSS/2MASS and our own measurements for different stars around each target (after any existing clear offset has been corrected). If a galaxy was observed on several nights in the same filter we added the uncertainties in quadrature to obtain a total $\\sigma_{zp}$ for that galaxy and filter. We have further compared the photometry of our field stars to values from the Pickles stellar library in both optical and NIR, and found no significant offsets.\\\\ \\section[]{Methods}\\protect\\label{methods} \\noindent The methods used in obtaining surface brightness and color profiles, structural parameters and other quantities of interest were presented in detail in \\M. Here we will provide a brief outline of the major steps but we refer the interested reader to \\M~for a more in--depth description of the procedures, the individual sources of uncertainty, motivation for the error composition, systematic errors consideration, etc. For the sake of brevity hereafter we will refer to $Ks$ as simply $K$ (except in the conclusions). \\subsection[]{Contour plots and RGB images}% \\noindent Contour plots were obtained with a combination of the python \\emph{astLib} package and the built--in \\emph{pylab} function \\emph{contour}. The isophotal bin is $0.5$ mag for all galaxies. To reduce the noise in the fainter isophotes the images were partially smoothed with the boxcar median filter. The RGB images for each galaxy were made with our own implementation of the~\\citet{2004PASP..116..133L} algorithm, where the same scaling and stretch factors were applied to all galaxies in order to facilitate direct comparison. The contour plots and RGB images are both oriented so that North is up and East is to the left. To illustrate the difference in color schemes between the \\emph{SDSS} and our own RGB image we show such a comparison in Figure~\\ref{sdssRGB} for two random galaxies from our sample. Figure~\\ref{datafig} contains the contour and RGB plots for each galaxy. The individual boxcar median filter width and the isophotal level at which it was applied is also indicated in this figure. \\subsection[]{Surface brightness profiles}% \\noindent We obtained isophotal and elliptical integration surface brightness profiles for the galaxies in the sample. In the former case we used a constant magnitude bin size of $0.5^{m}$ for all galaxies and all filters, and the deepest image (usually the $B$ band) to define the area of integration at each step. This area was then applied to the rest of the filters. In the case of elliptical integration the radial bin size is $1$ arcsec for all galaxies and all filters, and we again used the $B$ band to define the parameters of the integrating ellipse and applied those to the rest of the filters. In other words, the same physical area is sampled at each magnitude or radial bin in all filters. All foreground and background sources, except the target galaxy, were masked out prior to performing surface photometry on the images, where the mask size is usually a factor of $2.5$ larger than what is returned by \\emph{SExtractor}. Though the source detection and masking procedures are automatic, all masks were visually inspected and modified if it was deemed necessary. \\\\  \\noindent The elliptical integration errors include the zero point uncertainty $\\sigma_{zp}$, the uncertainty in the sky $\\sigma_{sky}$, and the uncertainty in the mean flux level, represented by the standard deviation of the mean flux in each elliptical ring, $\\sigma_{sdom}$. The isophotal errors are similarly obtained but of course exclude $\\sigma_{sdom}$. The details of the error estimation and the integration procedures are described in \\M. Figure~\\ref{datafig} shows the isophotal and elliptical surface brightness profiles as well as the resulting radial color profiles for all galaxies.  \\begin{figure} \\centering \\includegraphics[width=8cm,height=4.cm]{fig1.ps} \\includegraphics[width=8cm,height=4.cm]{fig2.ps} \\caption{\\textbf{UM461} (top) and \\textbf{UM500} (bottom) in SDSS (left) and our (right) RGB color schemes. } \\protect\\label{sdssRGB} \\end{figure} \\subsection[]{Integrated surface photometry}\\protect\\label{integsurfphot} \\noindent Similar to \\M, in Tables~\\ref{pa_holm_tbl} and~\\ref{totlumtbl} we present the parameters derived from the surface photometry, including position angle and ellipticity, the Holmberg radius $r_H$, and the total apparent and absolute magnitudes for each galaxy measured down to $r_H$. The error on the total luminosity was obtained by varying the position angle and ellipticity parameters by $\\pm 5^\\circ$, respectively $\\pm0.1$. In Table~\\ref{morph_tab} we summarize some general information for each target, such as oxygen--based metallicities and $H_\\beta$ equivalent widths, as well as the morphological class obtained either from the literature or through our own analysis of the morphology, where such classification was missing. We also split the underlying host galaxy in two regions, one between $\\mu_B\\sim24$--$26$ and one between $\\mu_B\\sim26$--$28$ mag arcsec${}^{-2}$, and calculate the total color over these regions (Table~\\ref{totclrtbl}). In the same table we also provide the color of the central region, from the center of integration down to $\\mu_B\\sim24$ mag arcsec${}^{-2}$, which contains contributions from both the star forming regions and the underlying host galaxy.\\\\  \\noindent Most of the host galaxies in this sample are well approximated by a disk, so it is meaningful to estimate the scale length $h_r$ and the central surface brightness $\\mu_0$ (Table~\\ref{scalelentbl}) for the sample, which we do in the same way as in \\M. The assumption of knowing the exact shape of the underlying host enables us to give an estimate of the burst luminosity, i.e. the excess light above the exponential disk (Table~\\ref{burstclrtbl}). The burst errors include the fitting errors of the exponential disk, scaled to units of the profile errors, and the zero point uncertainty. Thus they may be underestimated since there is no measure of the uncertainty in the exact flux level of the burst region. The latter is not included since the burst region is never explicitly defined in $2D$. All color measurements are always performed over identical physical ranges for all filters, taking the $B$ band as reference for defining the respective regions and then applying these regions to the rest of the filters. The errors of the colors are the composite of the individual errors in the two filters, which in turn contain contributions from all three relevant sources of uncertainty -- $\\sigma_{zp}$, $\\sigma_{sky}$, and $\\sigma_{sdom}$. The structural parameters errors, $\\sigma(h_r)$ and $\\sigma(\\mu_0)$, are the propagated errors of the fitted slope, and a composite of the fit error and zero point uncertainty, respectively. \\subsection[]{Asymmetry and concentration}\\protect\\label{casparam} \\noindent Table~\\ref{asymtbl} shows the individual minimum Petrosian asymmetry ($A_P$) for each galaxy. These are calculated following~\\citet[][]{2000ApJ...529..886C} as $A=\\frac{\\sum\\lvert I_0-I_\\phi\\rvert}{2\\sum\\lvert I_0\\rvert}$, where $\\phi=180$ degrees. The measurements were performed over the area included in the Petrosian radius $r[\\eta(0.2)]$, where all pixels below the corresponding flux level are masked out. We use the inverted $\\eta$, defined as the ratio between the local surface brightness at some radius and the average surface brightness inside that radius~\\citep[see][and references therein]{2000AJ....119.2645B}. The individual Petrosian radii are also presented in Table~\\ref{asymtbl}, since it can be informative to know how large the enclosed area is.\\\\  \\noindent Alternative measures of asymmetry are shown in Table~\\ref{tab:altasym}, namely the Holmberg ($A_H^\\prime$) and the dynamical ($A_{dyn}$) asymmetry. $A_H^\\prime$ is calculated over images smoothed by a boxcar average filter of $1\\times1$ kpc from the area enclosed by the Holmberg radius at $\\mu=26.5$ mag arcsec${}^{-2}$ in the optical and $R_{23}$ at $\\mu=23$ mag arcsec${}^{-2}$ in the NIR. $A_{dyn}$, the dynamical asymmetry, is also calculated over smoothed images, but all pixels brighter than $\\mu=25$, $\\mu=21$ mag arcsec${}^{-2}$ are set to the constant value of $25,~21$ mag arcsec${}^{-2}$ in the optical, respectively the NIR. This means that all star forming regions contribute nothing to the total asymmetry, allowing $A_{dyn}$ to give more weight to the faint dynamical structures. The faintest isophote in $A_{dyn}$ is $27$ mag arcsec${}^{-2}$ in the optical and $23$ mag arcsec${}^{-2}$ in the NIR.\\\\  \\noindent The concentration index~\\citep[e.g.][]{2000AJ....119.2645B} was calculated from $C=5\\times\\log{\\frac{r_{80\\%}}{r_{20\\%}}}$ where $r_{20}$,$r_{80}$ are the radii at $20\\%$, respectively $80\\%$ light over an area inside the $1.5\\times r[\\eta(0.2)]$ radius. These values are listed in Table~\\ref{tab:conc}. ",
        "conclusions": "\\noindent In Figure~\\ref{SEMs} we compare the $B-V$ vs $V-K$ colors, measured over different physical regions of the galaxies, with the predictions from stellar evolutionary models. The model tracks with nebular emission (left column) assume zero redshift and instant burst, and are based on the \\emph{Yggdrasil} spectral synthesis code \\citep{2011ApJ...740...13Z}, whereas the pure stellar population tracks (right column) are based on \\citet{2008A&A...482..883M} isochrones, are also at zero redshift, and have an exponential SFR decay of $1$ Gyr. The total colors (first row) for about half the sample are redwards of $V-K\\sim2$, having a low to moderate contribution from nebular emission coming from the star forming regions. The rest of the galaxies bluewards of $V-K\\sim2$ have total $V-K$ colors clearly incompatible with those of a pure stellar population and may suffer the effects of nebular emission contamination and dust extinction. Examining only the burst color estimate down to $\\mu_B=24$ mag arcsec${}^{-2}$ we see that they are better fitted with nebular emission contribution, as expected. Most data points fall on or close to the model tracks, though some clearly are incompatible with this instant burst model, and would be better modeled with an extended burst instead. The fit can only get better when we account for the difference in used filters and the non-zero redshift of the sample. The very red $V-K$ outlier in that plot (Figure~\\ref{SEMs}, left column, second row) is UM422, or rather the composite galaxy together with the very extended neighbor. There are no visible blue knots of star formation inside the region $\\mu_B\\lesssim24$ mag arcsec${}^{-2}$ for that galaxy, so we expect the burst estimate to be very inaccurate for this target due to its morphology and profile shape, and the fact that we are not actually measuring a star forming region. The region $24\\lesssim\\mu_B\\lesssim26$ mag arcsec${}^{-2}$ shows colors similar to the total colors in the first row of Figure~\\ref{SEMs}, which implies that either the starburst is not dominating the total galaxy colors, or that the $24\\lesssim\\mu_B\\lesssim26$ region is still very much strongly contaminated by the burst contribution. The latter is not likely since we observe continuous smooth profile slopes along the entire $24\\lesssim\\mu_B\\lesssim28$ region, and also because we have established that the relative burst contribution is on average moderate. Further, from the predominantly flat color profiles we observe in Figure~\\ref{datafig} we would not expect a big change in the integrated color at larger radii. This is unfortunately difficult to see, since the $26\\lesssim\\mu_B\\lesssim28$ mag arcsec${}^{-2}$ region contains fewer measurements due to the limiting effect of the $K$ band. For the same reason, the errorbars here are large. Nevertheless, the location of the data points is here better fitted with the pure stellar population model, with some very metal--poor ($Z\\sim0.001$) and some very metal--rich $Z\\sim0.02$ hosts older than a few Gyr. The three outliers close to the solar--metallicity track beyond $V-K\\sim3$ are UM446, UM465, and UM499. Solar metallicity is not unusual for spiral galaxies, hence UM499 is not truly a deviating data point. With the help of $H\\alpha$ data in~\\citet{Paper3} we will be able to discern whether there is any nebular emission contribution at such radii for these three galaxies, and hence whether one should compare them to the left or right panel tracks in Figure~\\ref{SEMs}. For now we can conclude that the case for a so called ``red halo''~\\citep{2002A&A...390..891B,2005mmgf.conf..355B,2006ApJ...650..812Z} in this sample of emission line galaxies is weak, though if taken at face value the stellar evolutionary model tracks imply unusually high metallicities for some hosts in the sample. \\\\ \\begin{table*} \\begin{minipage}{150mm} \\caption{$A_{P}$ asymmetry in each filter measured inside the Petrosian radius $r[\\eta(0.2)]$, given here in kpc. In the optical the $S/N$ within the Petrosian radius is significantly high so that the typical errors are $\\ll0.02$ (rms), while in the NIR the $S/N$ often drops to low values of $\\lesssim400$ within $r[\\eta(0.2)]$, which gives typical errors of $\\sim0.05$. These typical errors have been estimated by~\\citet{2000ApJ...529..886C}.} \\protect\\label{asymtbl} \\begin{tabular}{|lccccccccccccccc|} \\hline Galaxy&$A_U$&$r_U$&$A_B$&$r_B$&$A_V$&$r_V$&$A_R$&$r_R$&$A_I$&$r_I$&$A_H$&$r_H$&$A_K$&$r_K$\\\\\\hline UM422&$0.58$&$7.7$&$0.43$&$7.7$&$0.44$&$7.6$&$0.39$&$7.2$&$0.40$&$7.4$&$0.43$&$7.2$&$0.55$&$5.4$\\\\ UM439&$0.26$&$1.2$&$0.26$&$1.3$&$0.27$&$1.3$&$0.22$&$1.5$&$0.21$&$1.7$&&&$0.28$&$1.6$\\\\ UM446&&&$0.11$&$0.7$&$0.11$&$0.7$&$0.13$&$0.9$&$0.11$&$1.0$&$0.40$&$1.0$&$0.32$&$1.0$\\\\ UM452&$0.27$&$1.5$&$0.24$&$1.7$&$0.23$&$1.8$&$0.24$&$1.8$&$0.23$&$2.0$&$0.31$&$2.0$&$0.27$&$2.0$\\\\ UM456&$0.36$&$1.9$&$0.39$&$2.1$&$0.43$&$2.4$&$0.40$&$3.0$&$0.35$&$3.4$&$0.55$&$2.4$&$0.50$&$2.6$\\\\ UM461&$0.45$&$0.7$&$0.42$&$0.7$&$0.47$&$0.7$&$0.44$&$0.7$&$0.35$&$0.9$&$0.48$&$0.9$&$0.49$&$0.8$\\\\ UM462&$0.16$&$0.9$&$0.22$&$0.9$&$0.23$&$0.9$&$0.23$&$1.1$&$0.28$&$1.1$&$0.42$&$1.2$&$0.34$&$1.2$\\\\ UM463&$0.17$&$0.5$&$0.15$&$0.5$&$0.20$&$0.5$&&&$0.13$&$0.6$&$0.35$&$0.6$&$0.37$&$0.7$\\\\ UM465&$0.19$&$0.3$&$0.09$&$0.6$&$0.16$&$0.8$&&&&&&&$0.14$&$1.2$\\\\ UM477&$0.37$&$0.6$&$0.22$&$9.3$&$0.19$&$8.0$&$0.21$&$7.4$&$0.15$&$7.3$&&&$0.30$&$0.8$\\\\ UM483&$0.17$&$1.1$&$0.14$&$1.2$&$0.12$&$1.2$&$0.12$&$1.2$&$0.13$&$1.2$&&&$0.15$&$1.2$\\\\ UM491&$0.23$&$1.1$&$0.22$&$1.2$&$0.21$&$1.2$&$0.20$&$1.4$&$0.19$&$1.4$&&&$0.20$&$1.4$\\\\ UM499&$0.20$&$1.7$&$0.18$&$2.2$&$0.14$&$6.1$&&&$0.08$&$6.3$&&&$0.09$&$3.5$\\\\ UM500&$0.63$&$4.7$&$0.36$&$4.7$&$0.42$&$4.7$&$0.41$&$4.7$&$0.31$&$4.9$&$0.45$&$4.7$&$0.42$&$4.7$\\\\ UM501&$0.49$&$3.0$&$0.46$&$3.3$&$0.43$&$3.0$&$0.44$&$3.8$&$0.47$&$4.1$&$0.65$&$2.4$&$0.59$&$3.5$\\\\ UM504&$0.20$&$0.6$&$0.13$&$0.7$&$0.14$&$0.7$&$0.11$&$0.8$&$0.10$&$0.8$&$0.30$&$0.8$&$0.13$&$0.8$\\\\ UM523A&$0.48$&$2.6$&$0.33$&$3.5$&$0.29$&$3.5$&$0.27$&$3.6$&$0.26$&$3.5$&&&$0.40$&$3.5$\\\\ UM523B&$0.30$&$1.7$&$0.18$&$1.7$&$0.17$&$1.7$&$0.15$&$1.7$&$0.15$&$1.7$&&&$0.24$&$1.8$\\\\ UM533&$0.41$&$1.5$&$0.22$&$2.1$&$0.21$&$2.4$&$0.22$&$2.6$&$0.16$&$2.9$&&&$0.41$&$2.6$\\\\ UM538&$0.28$&$0.5$&$0.23$&$0.7$&$0.22$&$0.8$&$0.19$&$0.9$&$0.19$&$0.9$&&&$0.37$&$0.8$\\\\ UM559&$0.31$&$2.5$&$0.17$&$2.4$&$0.15$&$2.4$&$0.23$&$2.5$&$0.34$&$2.5$&$0.60$&$2.7$&$0.33$&$2.4$\\\\\\hline \\hline \\end{tabular} \\end{minipage} \\end{table*} \\subsection{Asymmetries}\\protect\\label{asymdiscuss} \\noindent In Table~\\ref{asymtbl} and Figure~\\ref{asymhist} (left panel) we present the distribution of the Petrosian $A_{P}$ asymmetries measured in each filter down to the Petrosian radius $r[\\eta(0.2)]$. Since the sample consists exclusively of emission line galaxies, the composite asymmetry of a galaxy is usually dominated by the flocculent asymmetry, where we use the distinction ``flocculent'' and ``dynamical'' asymmetry as defined by~\\citet{2000ApJ...529..886C}. This domination is further enhanced by our choice of the area over which the asymmetry is measured -- Table~\\ref{asymtbl} shows that the Petrosian $r[\\eta(0.2)]$ radius is usually quite small, and hence the enclosed area is limited to fairly bright surface brightness levels. This implies that for many galaxies the dominating component to the total asymmetry (flocculent plus dynamical) will be the asymmetry due to individual star forming knots, i.e. the flocculent asymmetry, best estimated by $A_P$. The dynamical asymmetry contribution is underestimated in this way, because the tidal tails and plumes usually associated with mergers and strong tidal interactions, can be very faint and will thus lie beyond the $r[\\eta(0.2)]$ radius. Even if faint features were included their contribution would be negligible since the asymmetry is luminosity weighted. In an attempt to obtain a better estimate of the dynamical asymmetry component in \\M~we examined alternative asymmetry measurements, such as the ``Holmberg'' $A_H^\\prime$ asymmetry and the purely dynamical $A_{dyn}$ asymmetry. We will come back to these later on but first let us examine in detail the behavior of the Petrosian $A_P$ asymmetry.\\\\ \\begin{figure*} \\begin{center} \\includegraphics[width=18.0cm]{fig33.ps} \\caption{$B-V$ total color (Table~\\ref{totlumtbl}) vs. the $A_{P}$ asymmetry in the $R$ (left panel), and the $I$ bands (right panel). The markers are color--coded by their $B$ band asymmetry, while their size reflects their Holmberg radius $r_H$. The gray area is the location of the fiducial galaxy color--asymmetry sequence as defined in~\\citet{2000ApJ...529..886C}.}\\protect\\label{conselice} \\end{center} \\end{figure*} \\begin{table} \\caption{Minimum asymmetries measured in each filter. The two numbers per filter per galaxy are Holmberg $A_H^\\prime$ asymmetry measured over the area enclosed by the Holmberg radius $r(\\mu=26.5)$ in the optical and by $r(\\mu=23)$ in the NIR (top value), and the dynamical $A_{dyn}$ asymmetry, with regions $\\mu<25$ ($\\mu<21$) set to $25$ ($21$) mag arcsec${}^{-2}$ in the optical (NIR) (bottom value). The images are pre--processed by a boxcar average of size $1\\times1$ kpc.} \\protect\\label{tab:altasym} \\begin{tabular}{|l|r|r|r|r|r|r|r|} \\hline Galaxy&$A_U$&$A_B$&$A_V$&$A_R$&$A_I$&$A_H$&$A_K$\\\\ \\hline UM422 &$ 0.41$ &$ 0.36$ &$ 0.34$ &$ 0.30$ &$ 0.28$ &$ 0.25$ &$ 0.32$\\\\ &$ 0.18$ &$ 0.15$ &$ 0.15$ &$ 0.13$ &$ 0.12$ &$ 0.23$ &$ 0.29$\\\\ UM439 &$ 0.09$ &$ 0.09$ &$ 0.10$ &$ 0.15$ &$ 0.18$ &  &$ 0.10$\\\\ &$ 0.13$ &$ 0.12$ &$ 0.16$ &$ 0.21$ &$ 0.21$ &  &$ 0.11$\\\\ UM446 &  &$ 0.05$ &$ 0.05$ &$ 0.05$ &$ 0.06$ &$ 0.08$ &$ 0.10$\\\\ &  &$ 0.03$ &$ 0.05$ &$ 0.08$ &$ 0.11$ &$ 0.10$ &$ 0.12$\\\\ UM452 &$ 0.10$ &$ 0.11$ &$ 0.11$ &$ 0.11$ &$ 0.11$ &$ 0.21$ &$ 0.12$\\\\ &$ 0.10$ &$ 0.05$ &$ 0.08$ &$ 0.08$ &$ 0.10$ &$ 0.24$ &$ 0.11$\\\\ UM456 &$ 0.23$ &$ 0.28$ &$ 0.30$ &$ 0.28$ &$ 0.26$ &$ 0.35$ &$ 0.35$\\\\ &$ 0.11$ &$ 0.09$ &$ 0.08$ &$ 0.10$ &$ 0.13$ &$ 0.27$ &$ 0.32$\\\\ UM461 &$ 0.20$ &$ 0.20$ &$ 0.23$ &$ 0.20$ &$ 0.16$ &$ 0.09$ &$ 0.11$\\\\ &$ 0.09$ &$ 0.08$ &$ 0.08$ &$ 0.10$ &$ 0.10$ &$ 0.09$ &$ 0.10$\\\\ UM462 &$ 0.07$ &$ 0.09$ &$ 0.10$ &$ 0.10$ &$ 0.13$ &$ 0.10$ &$ 0.14$\\\\ &$ 0.10$ &$ 0.06$ &$ 0.08$ &$ 0.10$ &$ 0.12$ &$ 0.08$ &$ 0.08$\\\\ UM463 &$ 0.04$ &$ 0.02$ &$ 0.04$ &  &$ 0.09$ &$ 0.11$ &$ 0.08$\\\\ &$ 0.06$ &$ 0.02$ &$ 0.05$ &  &$ 0.06$ &$ 0.14$ &$ 0.06$\\\\ UM465 &$ 0.04$ &$ 0.04$ &$ 0.04$ &  &  &  &$ 0.04$\\\\ &$ 0.06$ &$ 0.05$ &$ 0.05$ &  &  &  &$ 0.05$\\\\ UM477 &$ 0.20$ &$ 0.14$ &$ 0.12$ &$ 0.11$ &$ 0.10$ &  &$ 0.08$\\\\ &$ 0.17$ &$ 0.11$ &$ 0.10$ &$ 0.12$ &$ 0.11$ &  &$ 0.09$\\\\ UM483 &$ 0.08$ &$ 0.07$ &$ 0.07$ &$ 0.08$ &$ 0.07$ &  &$ 0.08$\\\\ &$ 0.08$ &$ 0.07$ &$ 0.16$ &$ 0.12$ &$ 0.15$ &  &$ 0.07$\\\\ UM491 &$ 0.12$ &$ 0.12$ &$ 0.12$ &$ 0.11$ &$ 0.11$ &  &$ 0.10$\\\\ &$ 0.08$ &$ 0.06$ &$ 0.07$ &$ 0.11$ &$ 0.12$ &  &$ 0.06$\\\\ UM499 &$ 0.12$ &$ 0.09$ &$ 0.08$ &  &$ 0.07$ &  &$ 0.08$\\\\ &$ 0.15$ &$ 0.11$ &$ 0.18$ &  &$ 0.14$ &  &$ 0.14$\\\\ UM500 &$ 0.36$ &$ 0.28$ &$ 0.29$ &$ 0.27$ &$ 0.25$ &$ 0.25$ &$ 0.26$\\\\ &$ 0.11$ &$ 0.11$ &$ 0.11$ &$ 0.11$ &$ 0.12$ &$ 0.21$ &$ 0.22$\\\\ UM501 &$ 0.37$ &$ 0.34$ &$ 0.31$ &$ 0.30$ &$ 0.29$ &$ 0.25$ &$ 0.26$\\\\ &$ 0.13$ &$ 0.11$ &$ 0.11$ &$ 0.19$ &$ 0.24$ &$ 0.25$ &$ 0.26$\\\\ UM504 &$ 0.04$ &$ 0.05$ &$ 0.04$ &$ 0.05$ &$ 0.06$ &$ 0.06$ &$ 0.07$\\\\ &$ 0.08$ &$ 0.03$ &$ 0.03$ &$ 0.05$ &$ 0.19$ &$ 0.05$ &$ 0.06$\\\\ UM523A &$ 0.27$ &$ 0.24$ &$ 0.22$ &$ 0.21$ &$ 0.21$ &  &$ 0.22$\\\\ &$ 0.21$ &$ 0.10$ &$ 0.09$ &$ 0.13$ &$ 0.15$ &  &$ 0.13$\\\\ UM523B &$ 0.21$ &$ 0.10$ &$ 0.12$ &$ 0.10$ &$ 0.12$ &  &$ 0.10$\\\\ &$ 0.18$ &$ 0.14$ &$ 0.13$ &$ 0.13$ &$ 0.11$ &  &$ 0.12$\\\\ UM533 &$ 0.13$ &$ 0.08$ &$ 0.08$ &$ 0.08$ &$ 0.06$ &  &$ 0.09$\\\\ &$ 0.06$ &$ 0.03$ &$ 0.03$ &$ 0.06$ &$ 0.08$ &  &$ 0.06$\\\\ UM538 &$ 0.08$ &$ 0.06$ &$ 0.08$ &$ 0.06$ &$ 0.05$ &  &$ 0.04$\\\\ &$ 0.03$ &$ 0.04$ &$ 0.09$ &$ 0.07$ &$ 0.07$ &  &$ 0.04$\\\\ UM559 &$ 0.20$ &$ 0.12$ &$ 0.12$ &$ 0.11$ &$ 0.12$ &$ 0.21$ &$ 0.09$\\\\ &$ 0.16$ &$ 0.12$ &$ 0.12$ &$ 0.13$ &$ 0.20$ &$ 0.21$ &$ 0.09$\\\\ \\hline \\end{tabular} \\end{table} \\begin{table} \\caption{Concentration parameter for each filter} \\protect\\label{tab:conc} \\begin{tabular}{@{}|l|r|r|r|r|r|r|r|r|@{}} \\hline Galaxy&$C_U$&$C_B$&$C_V$&$C_R$&$C_I$&$C_H$&$C_K$\\\\\\hline UM422&$2.2$&$2.5$&$2.6$&$2.6$&$2.7$&$2.8$&$2.4$\\\\ UM439&$2.2$&$2.2$&$2.4$&$2.7$&$2.7$&&$2.7$\\\\ UM446&&$2.4$&$2.4$&$2.7$&$3.0$&$3.0$&$2.1$\\\\ UM452&$2.2$&$2.6$&$2.7$&$2.7$&$2.9$&$2.9$&$2.9$\\\\ UM456&$2.4$&$2.9$&$3.0$&$3.2$&$3.3$&$2.7$&$3.2$\\\\ UM461&$1.8$&$1.8$&$1.5$&$2.1$&$2.4$&$2.4$&$2.2$\\\\ UM462&$2.4$&$2.6$&$2.6$&$2.4$&$2.4$&$2.6$&$2.6$\\\\ UM463&$1.5$&$1.5$&$1.5$&&$2.0$&$2.0$&$2.0$\\\\ UM465&$1.5$&$2.1$&$2.8$&&&&$3.5$\\\\ UM477&$2.4$&$2.7$&$3.0$&$3.2$&$3.2$&&$3.3$\\\\ UM483&$1.8$&$2.1$&$2.1$&$2.1$&$2.1$&&$2.1$\\\\ UM491&$1.8$&$2.1$&$2.4$&$2.4$&$2.4$&&$2.4$\\\\ UM499&$2.4$&$2.6$&$4.0$&&$3.7$&&$3.2$\\\\ UM500&$2.4$&$2.6$&$2.4$&$2.7$&$2.4$&$2.6$&$2.3$\\\\ UM501&$3.0$&$2.7$&$2.7$&$2.9$&$2.7$&$2.5$&$2.8$\\\\ UM504&$1.5$&$2.0$&$2.4$&$2.7$&$2.7$&$2.7$&$2.7$\\\\ UM523A&$3.3$&$2.7$&$2.6$&$2.6$&$2.6$&&$2.5$\\\\ UM523B&$2.8$&$2.7$&$2.8$&$2.4$&$2.5$&&$2.5$\\\\ UM533&$2.4$&$2.8$&$2.8$&$2.9$&$2.8$&&$2.8$\\\\ UM538&$2.4$&$2.4$&$2.6$&$2.8$&$2.8$&&$2.6$\\\\ UM559&$2.2$&$2.1$&$2.4$&$2.5$&$2.5$&$2.7$&$2.5$\\\\ \\hline \\end{tabular} \\end{table}  \\noindent This is a volume limited sample with fairly similar redshifts (Table~\\ref{nedtable}). The difference in distance to the galaxies varies at most by a factor of $2$, and hence differences in measured asymmetry values are not due to simple resolution effects. All optical filters show a strongly peaked distribution at small $A_{P}\\sim0.2$, which we expect since the minimum asymmetry is usually found in the starbursting knots, regardless of their location in the galaxy. There are, of course, exceptions to this as is the case for galaxies with starforming regions of comparable brightness located symmetrically on either side of the geometric center, e.g. $UM559$, $UM483$ in this sample. For most galaxies, the physical location of the minimum asymmetry does not change with filter, however the value of the asymmetry does so, with the most drastic change observed when going from optical to NIR filters. The $H$ and $K$ histograms both show a distinct shift of the majority of the targets to much higher asymmetry values, $A_{P}\\sim0.4$, than in the optical. This change in asymmetry from optical to NIR reflects the change between flocculent to dynamical domination in the integrated asymmetry value. In the NIR the contribution to the light output of the galaxy from old stars is significant, while the burst is of diminished importance. Hence, we expect the NIR asymmetry to be a reflection of the departure from the symmetric ground state due to dynamical effects, such as merging or tidal interactions. We can distinguish three groups of galaxies based on their $A_{P}$ asymmetry behavior in the optical and NIR. \\begin{itemize} \\item \\textbf{Small optical and small NIR $A_{P}$ asymmetry:} These are predominantly members from the \\emph{nE} BCG class, have regular isophotes, and include all the nuclear starbursts. Thus, they have small flocculent and small dynamical asymmetries. They are $UM439$, $UM452$, $UM465$, $UM477$, $UM483$, $UM491$, $UM499$, and $UM504$. \\item \\textbf{Small optical and large NIR $A_{P}$ asymmetry:} This group is characterized by having a spatially extended burst region, or multiple SF knots off--center. The host is otherwise regular at faint isophotes, i.e. the dominant morphological classes here \\emph{iE} and \\emph{SS} BCGs. These are UM446, UM462, UM463, UM523A, UM523B, UM533, UM538, and UM559. This group has small flocculent and large dynamical asymmetries. \\item \\textbf{Large optical and large NIR $A_{P}$ asymmetry:} This group contains all galaxies with highly irregular morphologies and/or numerous SF knots. These are UM422, UM456, UM461, UM500, and UM501. Note that all targets classified as mergers (\\emph{iI}) are found here. This group has large flocculent and large dynamical asymmetries. \\end{itemize} \\noindent Interestingly, we found no galaxies with large optical and small NIR $A_P$ asymmetries in this sample. If we use optical asymmetry as a proxy for the flocculent asymmetry component, and NIR asymmetry as a proxy for the dynamical asymmetry component, then such a combination would imply a galaxy in or close to the symmetric ground state but with a non-nuclear starforming region. This is difficult to achieve because any compact localized off-center star formation presumably would require some sort of tidal interaction or merger to trigger, which in turn would raise the dynamical asymmetry value. While possible, such morphological setup is obviously rare. \\\\ \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{fig34.ps} \\caption{Petrosian asymmetry ($A_P$) vs. all alternative asymmetry measurements. $A_H$ is the Holmberg asymmetry without any smoothing of the images, while a $1\\times1$ kpc smoothing box has been applied to $A_P^\\prime$ (Petrosian), $A_H^\\prime$ (Holmberg), and dynamical $A_{dyn}$ asymmetry. The uninterrupted lines are least square fits to the similarly colored data points. The dashed line has a slope of $1$. The dotted line is an extrapolation from the correlation of~\\citet{2003ApJS..147....1C} for E, S0, Sa--b, Sc--d and Irr, which lie beyond $A_P\\sim0.6$. All asymmetries here are from the $I$ band, in order to facilitate comparison with~\\citet{2003ApJS..147....1C}.}\\protect\\label{allasymmetries} \\end{center} \\end{figure}  \\noindent There is a known correlation between the (blue) color and the (red) $A_P$ asymmetry for spheroids, disks and irregulars~\\citep{2000ApJ...529..886C}, and in Figure~\\ref{conselice} we compare the total $B-V$ color of the galaxies in the sample to the $R$ and $I$ asymmetries. We have estimated the region of the fiducial galaxy color-asymmetry sequence by~\\citet{2000ApJ...529..886C}, and plotted it for comparison. Galaxies deviating from this sequence are too asymmetric for their observed color, which is an indication of a merger or ongoing interaction. All of the potential mergers in the group with large optical and NIR asymmetries are indeed located to the right of the fiducial line. Star formation alone cannot account for their measured asymmetries, and a boost to the asymmetry by dynamical processes is needed. We note that our previous suspicions about the merger nature of UM500 are reinforced in this figure, with UM500 falling clearly to the far right of the fiducial line. We find this convincing and now firmly classify UM500 as a \\emph{iI,M} BCG. This means that the galaxy group with large optical and large NIR $A_P$ asymmetry now exclusively captures all classified mergers in our sample.\\\\  \\noindent The fact that the Petrosian $A_P$ asymmetry varies with wavelength regime is a strong indication that it is not completely flocculent dominated as was the case in \\M. If this is the case it should be also evident from the behavior of the alternative asymmetry measurements shown in Table~\\ref{tab:altasym} and described in \\S~\\ref{casparam}. The right panel of Figure~\\ref{asymhist} shows the distribution of the Holmberg $A_H^\\prime$ and the dynamical $A_{dyn}$ asymmetries. The $A_H^\\prime$ distribution differs significantly from the $A_P$ values. The grouping by optical/NIR asymmetry we presented above is now destroyed, and only two major groups emerge -- one with small ($\\lesssim0.2$) optical and small NIR $A_H^\\prime$ which contains the majority of the galaxies, and one with large ($\\gtrsim0.2$) optical and large NIR $A_H^\\prime$, which contains UM422, UM456, UM500, UM501, and UM523A. The fact that $A_H^\\prime$ does not follow the same distribution as $A_P$ implies that, contrary to \\M, the latter is not completely flocculent dominated and the contribution of the dynamical component to $A_P$ is not insignificant. We also note that the $K$ band $A_H^\\prime$ asymmetry is nearly identical to the dynamical $A_{dyn}$ asymmetry distribution. The same is true for the $H$ band but is better seen comparing the $A_H^\\prime$ and $A_{dyn}$ values in Table~\\ref{tab:altasym}. This confirms our assumption that the NIR asymmetry is a very good proxy of the dynamical asymmetry component. Further, in the $A_{dyn}$ asymmetry distribution there are essentially no galaxies with large ($\\gtrsim0.2$) optical asymmetries. This is consistent with the optical asymmetry being flocculent dominated, a dominance which is effectively neutralized through our method of obtaining $A_{dyn}$. In terms of morphological class we would expect $A_{dyn}$ to only be able to distinguish mergers from the rest. Indeed, the galaxies here fall into two groups, the majority having small ($\\lesssim0.2$) optical and NIR $A_{dyn}$ asymmetries. The exceptions are UM422, UM456, UM500, and UM501, all of which are mergers, and have optical $A_{dyn}\\sim0.1$ and a much larger NIR $A_{dyn}\\sim0.3$ (averaged values). In other words, our measure of the purely dynamical component, $A_{dyn}$, successfully represents the effect mergers have on the morphology of a galaxy. Note that $A_H^\\prime$ performs nearly as well in distinguishing mergers from non-mergers (one false positive notwithstanding), and is therefore an acceptable measure of the dynamical asymmetry for the low luminosity BCGs in our sample.\\\\  \\noindent In Figure~\\ref{allasymmetries} we plot these alternative asymmetry measurements versus the Petrosian $A_P$ asymmetry. Similar to \\M, we find a strong correlation between the smoothed and unsmoothed Holmberg ($A_H^\\prime,~A_H$) and Petrosian ($A_P^\\prime,~A_P$) asymmetries. The correlation coefficients for $A_P$ vs $A_P^\\prime,~A_H^\\prime,~A_H$ are $R=0.8,~0.9,~0.9$, with line slopes $\\mathfrak{m}=0.60,~0.62,~0.87$ respectively. Contrary to \\M~we also find a medium strength correlation between $A_P$ (predominantly flocculent) and $A_{dyn}$ (dynamical) asymmetries. Pearson's $R$ for this correlation is $\\sim0.4$, which is only marginally significant for a sample of this size. This is consistent with our findings for the BCG sample of \\M.\\\\ \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{fig35.ps} \\caption{$G^B$ (blue circles) and $G^R$ (red circles) in a Holmberg radius ($r_H$) vs. dynamical $B$ band asymmetry ($A_{dyn}$) parameter space. The black/red circles are the deviating $G^R$ galaxies UM461 and UM500. The two spiral galaxies are not included in the plot.}\\protect\\label{holmdynA} \\end{center} \\end{figure}  \\noindent Figure~\\ref{holmdynA} shows the behavior of the blue $G^B$ and red $G^R$ BCGs (defined in \\S~\\ref{colortrends}) in terms of the dynamical asymmetry component $A_{dyn}$ and Holmberg radius, $r_H$. There is only a weak trend for the $G^B$ galaxies to have larger Holmberg radii and hence to be more extended but clearly they have higher dynamical asymmetries than the $G^R$ galaxies. While we established in \\S~\\ref{colortrends} that the star formation morphologies between blue and red BCGs clearly vary, the dynamical asymmetries suggest that also the morphologies of the underlying hosts of blue and red BCGs are different. As we already saw in the preceding discussion some of the $G^B$ galaxies are clearly mergers or show signs of strong tidal interaction. Other members of the $G^B$ group, e.g. UM491 or UM483, seem unlikely merger candidates yet they, too, display high dynamical asymmetries and blue colors. It is possible that the dynamical asymmetry component allows us to identify not only the obvious major mergers but also the more ``mature'' or minor ones. For example, the location and morphology of the star forming regions in UM491 or UM483 are suggestive of a long passed dynamical disturbance. We will investigate the ability of $A_{dyn}$ to detect minor mergers in detail in~\\citet{Paper3}. \\\\  \\noindent The galaxies in our sample are clearly separated from spheroids and early and late type disks in the concentration--asymmetry parameter space (Figure~\\ref{concVSasym}). They occupy the same region as the BCGs in \\M, with large asymmetries and small concentration indices. Note that BCGs show significant scatter in the concentration--asymmetry parameter space compared to normal galaxies."
    },
    "1208/1208.0598_arXiv.txt": {
        "abstract": "Discovered in October 2010 by the LINEAR survey, P/2010 TO20 LINEAR-Grauer (\\LG) was initially classified as an inert Jupiter Trojan.  Subsequent observations obtained in October 2011 revealed \\LG\\ to be a Jupiter-family comet. \\LG\\ has one of the largest perihelia ($q=5.1$ au) and lowest eccentricities ($e=0.09$) of the known JFCs. We report on observations of \\LG\\ taken on 29 October 2011 and numerical simulations of its orbital evolution. Analysis of our data reveals that \\LG\\ has a small nucleus ($<3$ km in radius) with broadband colours ($B-R=0.99\\pm0.06$ mag, $V-R=0.47\\pm0.06$ mag) typical of JFCs. We find a model dependent mass-loss rate close to 100 kg s$^{-1}$, most likely powered by water-ice sublimation. Our numerical simulation indicate that the orbit of \\LG\\ is unstable on very short (10 to 100 yr) timescales and suggest this object has recently evolved into its current location from a more distant, Centaur-type orbit. The orbit, dynamics and activity of \\LG\\ share similarities with the well known case of comet 29P/Schwassmann-Wachmann 1. ",
        "introduction": "Jupiter family comets (JFCs) originate in the transneptunian region of the solar system known as the Kuiper belt. Kuiper belt objects (KBOs) preserve key information about the epoch of planetesimal formation. The low ambient temperature ($\\sim40$ K) allows KBOs to retain most if not all of the ices present in their formation environment. This makes them valuable time capsules with which to study an epoch long gone. JFCs represent the small end of the steep size distribution of KBOs. They are survivors of a dynamically intermediate population, the Centaurs, that were neither ejected from the solar system nor collided with one of the giant planets (Jupiter to Neptune) on their journey from the Kuiper belt into the inner solar system.  The surfaces of JFCs are heavily processed when compared to KBOs, mainly by sublimation of surface ice. Indeed, the optical broadband colours of the two populations differ substantially: JFCs have typically solar (neutral) colours ($B-R\\sim1.3$, $g-r\\sim0.6$), with a spread of about 0.2 mag \\citep{2009Icar..201..674Lamy, 2012Icar..218..571Solontoi} while KBOs span a broad range of optical colours, from neutral ($B-R\\sim1$) to very red \\citep[$B-R\\sim2.5$,][]{2001AJ....122.2099J, 2002AJ....123.1039Jew}. The neutral surfaces are usually attributed to a fresh ice coating while very red surfaces are thought to be the end product of irradiation of initially neutral material \\citep{1987JGR....9214933Thompson, 1998Icar..134..253Moroz}.  The intermediate Centaurs present a peculiar, also intermediate distribution of colours which includes neutral objects (blue group) and very red objects (red group) but very few cases in between \\citep{2003A&A...410L..29Pei}. This Centaur colour bimodality may be related to their dynamical history \\citep{2012A&A...539A.144Melita} or simply to the fact that they are small \\citep{2012arXiv1206.3153Peixinho}.  Some objects straddle the line between Centaurs and JFCs. They are active, like JFCs, but have Centaur-like orbits. A famous example is 29P/Schwassmann-Wachmann 1 (hereafter 29P). The orbit of 29P is nearly circular ($e=0.04$) with perihelion and semimajor axis beyond Jupiter ($q=5.7$ au and $a=6.0$ au), and obeys the dynamical definition of a Centaur.  29P is constantly active and displays sporadic photometric outbursts \\citep{1958PASP...70..272Roemer, 1990ApJ...351..277Jewitt, 2008A&A...485..599TrigoRodriguez}.  The activity of 29P is mainly driven by carbon monoxide sublimation \\citep{1994Natur.371..229Senay, 1995Icar..115..213Crovisier}. Comet 29P has a radius $r=23\\pm3$ km (Yan Fern\\'andez, private comm.).  In this paper we present a study of comet P/2010~TO20 LINEAR-Grauer (hereafter \\LG) which is in many respects similar to 29P. Comet \\LG\\ has a low eccentricity ($e=0.09$), high perihelion ($q=5.1$ au) orbit, very close to meeting the Centaur definition (see Table \\ref{Table.Elements}, Fig.\\ \\ref{Fig.Orbit}). When \\LG\\ was discovered by the Lincoln Near Earth Asteroid Research (LINEAR) survey, on 1 October 2010, it was misclassified as a Trojan and its weak activity went unnoticed. We present optical broadband measurements of \\LG\\ and numerical simulations of its orbital evolution. The former show that \\LG\\ is weakly active, at a level comparable to the active Centaurs \\citep{2009AJ....137.4296Jewitt} while the latter indicate that \\LG\\ has in all likelihood recently evolved into its current location from a Centaur orbit. The implication is that \\LG\\ may offer a glimpse of a relatively small and fresh Centaur object.  \\begin{figure} \\includegraphics[width=83mm]{Orbits.eps} \\centering  \\caption{Osculating orbit of \\LG\\ on 3 October 2010 (near the time it was discovered by LINEAR) and on 29 October 2011, when the observations reported here were taken. Due to the proximity to Jupiter the osculating orbit of \\LG\\ changed significantly in a period of only one year. An arrow marks the direction to the last perihelion, the orbits of Earth, Mars, Jupiter and Saturn are plotted as dashed lines and the axes are labelled in astronomical units.}  \\label{Fig.Orbit} \\end{figure}  \\begin{figure*} \\includegraphics[width=140mm]{Stacks.eps} \\centering  \\caption{Summed stacks of 9 R-band (left) and 6 B-band (right) images of comet \\LG. The projected tail is more than 1\\arcmin\\ long and points in the direction opposite the orbital motion ($-v$).  Because of the very small solar phase angle (Table \\ref{Table.Geometry}), the antisolar direction lies nearly perpendicularly to the plane of the sky. For this reason, the near-nucleus tail has only a small component in the projected antisolar direction ($-\\sun$; see Fig.\\ \\ref{Fig.FilteredInsets}) which rapidly bends towards the direction opposite the orbital motion.  }  \\label{Fig.Stacks} \\end{figure*}  \\begin{figure*} \\includegraphics[width=140mm]{FilteredInsets.eps} \\centering  \\caption{Enlarged versions of the near-nucleus tail structure from the R-band stack shown in Fig.\\ \\ref{Fig.Stacks}.  The original image (left) is shown together with versions processed using Laplacian (centre) and median (right) filters. A trail of larger particles leaving the nucleus and trailing the comet's orbital motion is visible in the Laplacian-filtered image. The median-filtered image highlights the near-nucleus part of the tail which points in the projected antisolar direction. }  \\label{Fig.FilteredInsets} \\end{figure*}  \\begin{table} \\caption{Orbital properties of LINEAR-Grauer on 2011 Oct 29. \\label{Table.Elements}} \\begin{tabular}{@{}lr@{}} \\hline Property & Value \\\\ \\hline Semimajor axis, $a$                    & 5.610 au \\\\ Eccentricity, $e$                      & 0.087  \\\\ Inclination, $i$                       & 2$\\fdg$628  \\\\ Argument of perihelion, $\\omega$       & 252$\\fdg$9  \\\\ Longitude of ascending node, $\\Omega$  & 44$\\fdg$03  \\\\ Mean anomaly, $M$                      & 84$\\fdg$22  \\\\ True anomaly, $\\nu$                    & 94$\\fdg$19 \\\\ Last perihelion passage                & 2008 Sep 5 \\\\ Perihelion distance, $q$               & 5.122 au \\\\ Aphelion distance, $Q$                 & 6.097 au \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table} \\caption{Observing Geometry. \\label{Table.Geometry}} \\begin{tabular}{@{}lrr@{}} \\hline Property & 2010 October 03 & 2011 October 29 \\\\ \\hline Solar phase angle ($\\alpha$)            & 2.07\\degr & 0.89\\degr\\\\ Heliocentric distance ($R$)             & 5.295 au  & 5.603 au \\\\ Geocentric distance ($\\Delta$)          & 4.309 au  & 4.613 au \\\\ 1\\arcsec\\ at distance $\\Delta$          & 3135 km   & 3355 km  \\\\ \\hline \\end{tabular} \\end{table}  \\begin{table*} \\centering \\begin{minipage}{160mm} \\caption{Photometry Regions \\label{Table.PhotometryRegions}} \\begin{tabular}{@{}ccccccc@{}} \\hline Region & Region & Aperture radius & Aperture radius  & Projected radius & Magnitudes & Dominant Source \\\\ Label  & Shape  & [pixels] & [\\arcsec] & [km] & in Region & \\\\ \\hline $\\mathcal{R}_{4}$     & Circle  & $r=4$        & $\\phi=0.96$          & $d=3,220$            & $m_4$      & Nucleus (+ Coma) \\\\ $\\mathcal{R}_{4,8}$   & Annulus & $r=4$ to 8   & $\\phi=0.96$ to 1.92  & $d=3,220$ to 6,440   & $m_{4,8}$  & Nucleus + Coma   \\\\ $\\mathcal{R}_{8,13}$  & Annulus & $r=8$ to 13  & $\\phi=1.92$ to 3.12  & $d=6,440$ to 10,465  & $m_{8,13}$  & Coma            \\\\ $\\mathcal{R}_{13,20}$ & Annulus & $r=13$ to 20 & $\\phi=3.12$ to 4.80  & $d=10,465$ to 16,100 & $m_{13,20}$ & Coma      \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}   \\section[]{Observations} \\label{Sec.Observations}  \\LG\\ was observed on 29 October 2011 at the ESO New Technology Telescope (NTT) located at the La Silla Observatory, Chile.  The night was photometric and the seeing varied between 0.8 and 1.0\\arcsec. At the NTT we used the EFOSC2 instrument \\citep{1984Msngr..38....9Buzzoni, 2008Msngr.132...18Snodgrass} which is installed at the f/11 Nasmyth focus and is equipped with a LORAL $2048\\times2048$ CCD. We used the $2\\times2$ binning mode to bring the effective pixel scale to 0.24\\arcsec/pixel. Our observations were taken through Bessel $B,V,R$ filters (ESO \\#639, \\#641, \\#642, respectively).  The images of \\LG\\ were collected in a relatively regular fashion, in sets of three consecutive exposures per filter, with exposure times of 240 s for $B$ and 120 s for $V$ and $R$. In total, we collected 12 images in the $B$ band, 12 in $V$, and 18 in $R$. Throughout the observations of \\LG, the telescope was set to track the non-sidereal motion of \\LG\\ at an approximate rate of $-18$\\arcsec/hour in right ascension and $-6$\\arcsec/hour in declination.  Bias calibration frames and dithered twilight flats through all three filters were collected on the same night as the science data. The reduction of the science images, consisting of standard bias subtraction and flat fielding, was done using the IRAF {\\tt ccdproc} routines.  The R band images suffered from fringing which was removed using an IRAF package optimised for EFOSC2 that was kindly supplied by Colin Snodgrass. The \\LG\\ data were absolutely calibrated using observations of \\citet{1992AJ....104..340Lan} stars taken throughout the night.  \\LG\\ was observed as part of the Pan-STARRS PS1 all-sky survey very near the time of its discovery by LINEAR. Located on Haleakala, Maui, the 1.8-m PS1 telescope is equipped with a 1.4 gigapixel camera covering $3.2\\degr\\times3.2\\degr$ on the sky. The survey repeatedly covers the $3\\pi$ steradians of sky visible from Haleakala.  PS1 uses a photometric system that approaches the SDSS filter system with the addition of a wide ($w$) band filter which roughly corresponds to the combined pass band of the $gri$ filters \\citep{2012ApJ...750...99Tonry}. The PS1 gigapixel camera has a pixel scale 0.25\\arcsec/pixel.  \\LG\\ was imaged in four consecutive 45 s exposures taken through the $w$ filter on 3 October 2010. All images were processed automatically by the Pan-STARRS Image Processing Pipeline. ",
        "conclusions": "We report photometric observations and numerical simulations of the orbital evolution of the unusual comet P/2010 TO20 LINEAR-Grauer (\\LG). Our findings can be summarised as follows:  \\begin{enumerate}  \\item Comet \\LG\\ was active at the time of discovery (October 2010) by LINEAR, and remains active in October 2011. LINEAR did not detect the activity and initially misclassified \\LG\\ as a Trojan, which suggests that several active objects may have gone undetected by the survey.  \\item The nucleus of \\LG\\ has equivalent radius $r_e<3$ km, and colours $B-R=0.99\\pm0.06$ mag and $V-R=0.47\\pm0.06$ mag, values typical of Jupiter family comets. The data suggest a slight reddening of the dust colour with distance from the nucleus but the uncertainties are large.  We find no significant rotational photometric variability from the nucleus region.  \\item We obtain a model-dependent estimate of the mass-loss rate from \\LG\\ of $\\sim$100 kg s$^{-1}$. We favour water-ice sublimation as the simplest and most likely cause for activity in comet \\LG.  \\item Our numerical simulations show that the orbit of \\LG\\ is unstable on very short timescales and suggest that it may be a Centaur that recently arrived in the inner solar system, although other possibilities exist involving non-gravitational effects.  \\item Comet \\LG\\ is in a number of ways reminiscent of the well-known 29P/Schwassmann-Wachmann 1. 29P is a comet/active Centaur that shows sporadic outbursts superimposed on a background of constant activity. Comet \\LG\\ is an order of magnitude smaller (three orders of magnitude less massive) than 29P and yet displays similar activity per unit area. Comets 29P and \\LG\\ are interesting as possible examples of relatively unprocessed Centaurs.  \\end{enumerate}"
    },
    "1208/1208.4658_arXiv.txt": {
        "abstract": "Together, the Fermi-LAT and HESS have revealed the presence of an unusual GeV-TeV source coincident with Sgr A* at the Galactic center. Its high-energy emission appears to be bimodal, hinting at an energizing process more sophisticated than mere shock acceleration. It has been suggested that this may be evidence of strong, rapid variability, as required if Sgr A*'s emission were responsible for the fluorescent X-ray echos detected in nearby molecular clouds. In this {\\it Letter}, however, we show that stochastic acceleration in a more realistic two-phase environment surrounding the central black hole can accommodate Sgr A*'s high-energy spectrum quite well. We therefore suggest that the Fermi-HESS data alone do not necessarily provide evidence for strong variability in Sgr A*. ",
        "introduction": "The first 25 months of observations of the Galactic center region by the Fermi-LAT have revealed the presence of a source, 1FGL J1745.6-2900, at energies above 10 GeV, coincident with the TeV central object HESS J1745-290 (Chernyakova et al.~2011). The combined Fermi-HESS spectrum appears to be inflected, with a relatively steep region at intermediate energies, connected to flatter components at both low and high energies. More specifically, although the spectrum in the 100 MeV--300 GeV energy range can be fit by a power law with slope $\\alpha=2.212\\pm0.005$ and a flux normalization $F(100\\,{\\rm MeV})=(1.39\\pm0.02)\\times 10^{-8}$ cm$^{-2}$ s$^{-1}$ MeV$^{-1}$, a significantly better fit is obtained with a power-law index $\\alpha=2.196 \\pm0.001$ between 300 MeV and 5 GeV, and a separate power law with $\\alpha=2.681\\pm0.003$ in the range 5 GeV--100 GeV. The latter is also significantly steeper than the spectrum reported by the HESS collaboration, where $\\alpha\\approx 2.1$, with a total flux $1.87\\pm 0.30) \\times 10^{-8}$ m$^{-2}$ s$^{-1}$ above 1 TeV (Aharonian et al. 2009).  The Fermi-HESS source appears to be the high-energy counterpart to Sgr A* (see Melia 2007 for a review). We note, however, that although the HESS source is coincident within $\\sim30^{\\prime\\prime}$ of Sgr A*, its centroid is displaced roughly $7^{\\prime\\prime}$ ($\\sim 0.4$ pc) to the east of the Galactic center, an issue that we shall revist below.  In earlier work (Liu et al. 2006; Ballantyne et al. 2007), we considered the interesting possibility that the TeV $\\gamma$-rays are produced via $\\pi^0$ decays generated when relativistic protons, accelerated near the black hole, collide with hadrons farther out. Such a scenario is intriguing because our basic understanding of Sgr A* and its nearby environment precludes any possibility of this object producing a significant flux of TeV photons directly (Liu \\& Melia 2001). However, protons can be energized to TeV energies by stochastic acceleration in a magnetically dominated funnel within 20--30 Schwarzschild radii of the event horizon. Then, as these cosmic rays diffuse out through the surrounding medium, they may scatter with hydrogen nuclei in a shocked stellar wind region and with the circumnuclear molecular torus surrounding Sgr A*.  But the latest $\\sim 10$ GeV Fermi-LAT observations, and more recent theoretical work with diffusive particle acceleration (see, e.g., Fatuzzo et al. 2010; Melia \\& Fatuzzo 2011; Fatuzzo \\& Melia 2012), render this basic picture inadequate for several reasons. First, the combined Fermi-HESS measurements show that a simple extrapolation of the TeV spectrum into the GeV range does not provide an acceptable explanation for the data. Second, it now appears unlikely that the energy of a cosmic ray remains constant as it diffuses out to larger radii. Thus, our earlier conclusion---that the incipient proton power-law index has to be $\\approx0.75$ near the black hole to produce the observed photon index $\\approx 2.25$ farther out---is not valid when the cosmic rays undergo continued acceleration on their way out into the circumnuclear environment.  Attempting to explain the surprising Fermi-HESS spectrum of the source 1FGL J1745.6-2900, Chernyakova et al. (2011) modified the Ballantyne et al. (2007) treatment in several ways. Firstl, instead of adopting a two-phase medium surrounding Sgr A*, in which the relatively cold inner molecular torus is surrounded by a hotter, more tenuous wind-shocked region, they assumed a single plasma with a $1/r^2$ density profile. Secondly, they invoked a time-dependent cosmic-ray injection process which, by virtue of diffusion-induced and energy-dependent time delays, can produce the kind of inflected spectrum seen by the Fermi-LAT and HESS. They were able to produce a reasonable fit to the data, but only under the assumption that Sgr A* is highly variable on timescales of several hundred to several thousand years.  However, Chernyakova et al.'s (2011) analysis appears to be incomplete and inconclusive because several important aspects of the particle diffusion used by them are unrealistic. First and foremost, the assumption of a constant particle energy following ejection is difficult to justify, given what we know about the physical conditions surrounding Sgr A*. One of the principal goals of this {\\it Letter} is in fact to incorporate this important effect into the calculation of the spectrum.  Second, the supposition that protons are ejected episodically by Sgr A* is motivated by the detection of reflected X-ray emission from the Sgr B2 cloud complex some 100 parsecs away (and other similar clouds at various distances from the center). This Iron-line fluorescence is often interpreted as the light-echo of a flare produced by the supermassive black hole hundreds (or thousands) of years ago (see, e.g., Sunyaev et al. 1993, Fromerth, Melia, and Leahy 2001, Revnivtsev et al. 2004, Terrier et al. 2010). But this kind of phenomenon requires a variation in the overall power of Sgr A* by some 6 orders of magnitude during a very short time compared to its age. Though one cannot decidedly argue against such an unprecedented change, it is far more likely that the source of light responsible for these echoes was the impact onto the 50 km s$^{-1}$ molecular cloud behind Sgr A* by the supernova shell of the explosion that produced Sgr A East (see Fryer et al. 2006).  However, Chernyakova et al. (2011) also considered a steady-state proton ejection and were able to account for the Fermi data, albeit under unusual cirumstances. In order for the proton diffusion to produce the required stratified proton energy distribution, with the highest energy protons moving more or less rectilinearly towards larger radii, they found that the diffusion coefficient must be strongly dependent on the energy (i.e., $D(E)\\propto E^{0.65}$). A second principal goal of this {\\it Letter} is to calculate $D(E)$ from first principles, using the physical conditions prevalent around Sgr A*, to see if such an energy dependence is realistic. As we shall see, our results do not support such a strongly variable diffusion coefficient.  What this means, of course, is that it may be too simplistic to do away with the two-phase gaseous environment surrounding Sgr A*. The physical conditions in the circumnuclear disk are drastically different from those in the rest of the tenuous, hot medium filling the inner 3 parsecs. Adopting an average particle density, and a concomitantly simplified magnetic field structure in lieu of the actual variation of these quantities between the molecular and ionized phases, significantly alters the particle diffusion. For these three principal reasons, it is therefore necessary to examine whether a steady-state cosmic-ray ejection by Sgr A* can provide a reasonable alternative explanation for its 100 MeV--100 TeV spectrum, but without having to invoke a diffusion coefficient that depends strongly on the particle energy. ",
        "conclusions": "We have lent support to the idea that---given the right environment---cosmic rays can be accelerated to TeV energies via stochastic processes in a turbulent magnetic field. Although this has been suspected for many years, we can only now demonstrate this quantitatively by comparing detailed numerical simulations with high quality HESS and Fermi-LAT data. Earlier, we showed that the cosmic-ray population permeating the inner several hundred parsecs of the Galaxy could be accounted for with this process, but only with the physical conditions encountered in that region (Fatuzzo \\& Melia 2012). As such, the fact that such a relativistic particle distribution is found only near the Galactic center can be understood as a product of the unique conditions found there.  In this {\\it Letter}, we have focused on the particle acceleration occurring near Sgr A* itself. The combined data exhibited in Figure~1 present quite a challenge to any effort at understanding how these cosmic rays are produced. Related to this question is the issue of whether or not these observations provide evidence for strong variability in Sgr A*, as suggested by the fluorescent X-ray echos detected in nearby molecular clouds.  But here we have shown that the use of a molecular torus and an interstellar medium filled with shocked stellar winds can provide just the right framework for steady-state stochastic acceleration to produce the observed two-component spectrum. We therefore suggest that Sgr A*'s high-energy spectrum does not necessarily provide evidence for the kind of rapid variability required to produce the fluorescent X-ray echos."
    },
    "1208/1208.1761_arXiv.txt": {
        "abstract": "The stability properties of a low density ultra relativistic pair beam produced in the intergalactic medium by multi-TeV gamma-ray photons from blazars are analyzed.  The problem is relevant for probes of magnetic field in cosmic voids through gamma-ray observations. In addition, dissipation of such beams could affect considerably the thermal history of the intergalactic medium and structure formation.  We use a Monte Carlo method to quantify the properties of the blazar induced electromagnetic shower, in particular the bulk Lorentz factor and the angular spread of the pair beam generated by the shower, as a function of distance from the blazar itself.  We then use linear and nonlinear kinetic theory to study the stability of the pair beam against the growth of electrostatic plasma waves, employing the Monte Carlo results for our quantitative estimates.  We find that the fastest growing mode, like any perturbation mode with even a very modest component perpendicular to the beam direction, cannot be described in the reactive regime. Due to the effect of non-linear Landau damping, which suppresses the growth of plasma oscillations, the beam relaxation timescale is found significantly longer than the inverse Compton loss time.  Finally, density inhomogeneities associated with cosmic structure induce loss of resonance between the beam particles and plasma oscillations, strongly inhibiting their growth.  We conclude that relativistic pair beams produced by blazars in the intergalactic medium are stable on timescales long compared to the electromagnetic cascade's.  There appears to be little or no effect of pair-beams on the intergalactic medium. ",
        "introduction": "\\label{intro:sec} Streaming relativistic particles are common in tenuous astrophysical plasma and their propagation and stability properties a recurrent theme.  Examples include type III solar radio burst~\\citep{Benz93}, quasars' jets~\\citep{LS87}, cosmic-rays streaming out of star forming galaxies~\\citep{MB11} and cosmic-ray transport in the intracluster medium~\\citep{Ensslin11}. Propagation and stability properties, related in particular to the exitation of plasma waves, is subject of attentive investigation as they can play a crucial role in the interpretation of observational data.  Ultra-relativistic beams of $e^+e^-$ pairs are also generated in the intergalactic medium (IGM) by very high energy gamma-rays from distant blazars, by way of photon-photon interactions with the extragalactic background light~\\cite[EBL,][]{GS67,Schlickeiser12}. While blazars' spectra, and in particular their multi-TeV cut-off features, have been studied in detail to constraint the EBL~\\cite[e.g.][]{Aharonian06}, recently multi-GeV and TeV blazars observations have also been used to constrain magnetic field in cosmic voids for the first time~\\citep{NV10,Tavecchio10}. In fact, for flat enough blazar's spectra, the electromagnetic cascade should produce an observable spectral bump at multi-GeV energies. The absence of such a bump in a number of observed blazars is ascribed to the presence of a sufficiently strong magnetic field, $B_\\void\\gtrsim 10^{-16}$G, to deflect the pairs in less then an inverse Compton length, $\\ell_{IC}\\simeq$ Mpc$\\,(E_\\pm/{\\rm TeV})^{-1}(1+z)^{-4}$, where $z$ is the cosmological redshift~\\citep{Plaga95,NS09}. When time variability of the blazars is taken into account the above lower limit is relaxed to a more conservative value of $B_\\void\\gtrsim 10^{-18}$G~\\citep{Dermer11,Taylor11}. The required filling factor of the magnetic field is about 60\\%~\\citep{Dolag11}. Other potential effects of a magnetic field in voids on the electromagnetic cascades have also been investigated, including extended emission around gamma-ray point-like sources (Aharonian et al. 1994; Neronov \\& Semikoz 2007; Dolag et al. 2009; Elyiv et al. 2009; Neronov et al. 2010a) and the delayed \u201cechoes\u201d of multi-TeV gamma-ray flares or gamma-ray bursts (Plaga 1995; Takahashi et al. 2008; Murase et al. 2008, 2009).  However, in principle the pair-beam is subject to various instabilities, in particular microscopic plasma instabilities of the two-stream family.  On this account,~\\cite{Broderick12} conclude that transverse modes of the two-stream instability act on much shorter timescales than inverse Compton scattering, effectively inhibiting the cascade and invalidating the above magnetic field measurements.  In addition, as a result of the beam's relaxation, substantial amount of energy would be deposited into the IGM, with dramatic consequences for its thermal history~\\citep{Chang12,Pfrommer12}.  In this paper, we reanalyze the stability of blazars induced ultra-relativistic pair beams.  In particular, we use a Monte Carlo model of the electromagnetic shower to quantify the beam properties at various distances from the blazar, and analyze the stability of the produced beam following the work of Breizman, Rytov and collaborators~\\cite[reviewed in][]{BR74,B90}.  We find that even for very modest perpendicular components of the wave-vector, the analysis of the instability requires a kinetic treatment. We thus estimate the max growth rate of the instability and find that for bright blazars (with equivalent isotropic gamma-ray luminosity of 10$^{45}$ erg s$^{-1}$) it is suppressed by Coulomb collisions at distances $D\\gtrsim$ 50 and 20 physical Mpc at redshift 0 and 3, respectively. Importantly, the growth rate of plasma oscillations is found to be severely suppressed by non-linear Landau damping, so that even at closer distances to the blazar the beam relaxation timescale remains considerably longer than the inverse Compton cooling time.  Finally, the resonance condition cannot be maintained in the presence of density inhomogeneities associated to cosmological structure formation, which also act to dramatically suppress the instability. Thus our findings support the magnetic field based interpretation of the gamma-ray observational results and rule out effects of blazars' beam on the thermal history of the IGM.  Broderick et al. did not consider the role of density inhomogeneities and concluded that non-linear Landau damping is unimportant, although they did not present a quantitative analysis of the process.  The rest of this paper is organized as follows. Sec.~\\ref{pbeam:sec} summarizes the physical properties of pair beams produced by blazars and present the results of the Monte Carlo model. The two-stream instability in both the reactive and kinetic regimes is discussed in Sec.~\\ref{inst:sec}, where the max growth rate of the instability is also given and compared to the collisional rate.  Nonlinear effects are considered in Sec.~\\ref{stabil:sec}, where the timescales for the beam relaxation is derived. Finally, Sec.~\\ref{sum:sec} briefly summarizes the results. ",
        "conclusions": "\\label{sum:sec} We considered the stability properties of a low density ultra relativistic pair beam produced in the intergalactic medium by multi-TeV gamma-ray photons from blazars.  The physical properties of the pair beam are determined through a Monte Carlo model of the electromagnetic cascade.  In summary we find that the combination of kinetic effects, non-linear Landau damping and density inhomogeneities appear to considerably stabilize blazars induced ultra-relativistic beams over the inverse Compton loss timescale, so that the electromagnetic cascade remains mostly unaffected by the beam instability. This implies that the lack of a bumpy feature at multi-GeV energies in the gamma-ray spectrum of distant blazars cannot be attributed such instabilities and can in principle be related to the presence of an intergalactic magnetic field. Finally, heating of the IGM by pair beams appears negligible."
    },
    "1208/1208.6377_arXiv.txt": {
        "abstract": "{ Stars twinkle because their light propagates through the atmosphere. The same phenomenon is expected when the light of remote stars crosses a Galactic -- disk or halo -- refractive medium such as a molecular cloud. We present the promising results of a test performed with the ESO-NTT and the perspectives. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.6306.txt": {
        "abstract": "Star forming galaxies are thought to dominate the sub-mJy radio population, but recent work has shown that low luminosity AGN can still make a significant contribution to the faint radio source population. Spectral indices are an important tool for understanding the emission mechanism of the faint radio sources. We have observed the extended Chandra Deep Field South at 5.5 GHz using a mosaic of 42 pointings with the Australia Telescope Compact Array (ATCA). Our image reaches an almost uniform sensitivity of $\\sim$12 $\\mu$Jy rms over 0.25 deg$^2$ with a restoring beam of 4.9 $\\times$ 2.0 arcsec, making it one of the deepest 6cm surveys to date. We present the 5.5 GHz catalogue and source counts from this field. We take advantage of the large amounts of ancillary data in this field to study the 1.4 to 5.5 GHz spectral indices of the sub-mJy population. For the full 5.5 GHz selected sample we find a flat median spectral index, $\\alpha_{\\rm med} = -0.40$, which is consistent with previous results. However, the spectral index appears to steepen at the faintest flux density levels ($S_{5.5 GHz} < 0.1$ mJy), where $\\alpha_{\\rm med} = -0.68$. We performed stacking analysis of the faint 1.4 GHz selected sample (40 $<$ $S_{1.4 GHz}$ $<$  200 $\\mu$Jy) and also find a steep average spectral index, $\\alpha = -0.8$, consistent with synchrotron emission. We find a weak trend of steepening spectral index with redshift. Several young AGN candidates are identified using spectral indices, suggesting Gigahertz Peaked Spectrum (GPS)  sources are as common in the mJy population as they are at Jy levels. ",
        "introduction": "One of the most fundamental issues in astrophysics is when and how stars, galaxies and black holes form, and how they evolved with cosmic time. The black hole mass and bulge relation \\citep{kormendy1995, magorrian1998} suggests there is a link between active galactic nuclei (AGN) and star formation, hence studying the interaction between stars, galaxies and AGN phenomena is crucial for understanding galaxy formation in the early universe and how those galaxies evolve to the objects we see today. Radio emission can be produced by both AGN and star-forming processes, and thus radio wavelengths provide a unique window to study the cosmic evolution of these two important processes.  Bright radio sources ($> 100$ mJy) are associated with AGN activity (e.g. \\citealp{condon1984, georgakakis1999}), but the Euclidean-normalized radio source counts flatten below about 1 mJy and this cannot be explained by a population of radio-loud AGNs. A non-evolving population of local ($z < 0.1$) low-luminosity radio galaxies \\cite{wall1986}, strongly evolving normal spirals \\citep{condon1984, condon1989}, and star forming galaxies \\citep{windhorst1985, rowan-robinson1993} have all been suggested to explain this new population. The faint radio source counts have been successfully modelled by star forming galaxies \\citep{seymour2004, huynh2005}. The most commonly accepted paradigm has been that the sub-mJy population is largely made up of starforming galaxies. However a growing number of studies are finding that low luminosity AGN, both radio-loud and radio-quiet, still make a significant contribution to the sub-mJy population \\citep{jarvis2004, huynh2008, seymour2008, smolcic2008, padovani2009, padovani2011}.  There is conflicting evidence on the nature and properties of faint radio sources, and in particular their spectral index properties and whether there is a flattening of the average spectral indices for faint radio sources (e.g. \\citealp{randall2012}). \\cite{prandoni2006} found sources with 1.4 GHz flux densities less than a few mJy had an average 1.4 GHz to 5 GHz spectral index flatter than that of brighter radio sources. \\cite{owen2009} found a flattening of the average 325 MHz to 1.4 GHz  spectral index below $S_{1.4GHz} < 10$ mJy, but spectral indices appeared to steepen again at the faintest flux densities. In the Lockman Hole, the deepest radio field to date, no flattening of the 610 MHz to 1.4 GHz spectral indices was observed for  $S_{1.4GHz} > 0.1$ mJy \\citep{ibar2009}.  Galaxies dominated by star formation processes are expected to have a spectral index of $\\alpha = -0.8$ ($S \\propto \\nu^\\alpha$, \\citealp{condon1992}), consistent with synchrotron emission from electrons accelerated by supernovae. Although thermal Bremsstrahlung (free-free) emission found in HII regions can have a flatter spectral index, a flat or inverted spectrum is usually attributed to the superposition of different self-absorbed components of varying sizes at the base of the radio jet of a radio-loud AGN. A flattening in the average spectral index would imply there is a population of sources at sub-mJy flux densities with flat or inverted spectra, which are likely to be AGN.  Studying the spectral index properties of the faint radio population is important for understanding the $z-\\alpha$ relation, which is used to identify the highest redshift radio sources (e.g. \\citealp{debreuck2004, klamer2006}), identifying young radio AGN such as Gigahertz Peaked Spectrum sources and Compact Steep Spectrum sources \\citep{odea1998, randall2011}, and determining whether star formation or AGN processes are responsible for the radio emission in these sources.  Here we present 5.5 GHz observations by the Australia Telescope Large Area Survey (ATLAS, \\citealp{norris2006}) team of the extended Chandra Deep Field South which are well-matched in area to the extensive multiwavelength data in the region. This survey is important as it is the largest 6cm survey deeper than 100 $\\mu$Jy. We present the 5.5 GHz source catalogue and counts, and discuss the spectral index properties of the ATLAS radio sources. The paper is organized as follows: In Section 2 we describe the observations and data reduction. We discuss the source extraction and present the source catalogue in Section 3. Source counts are derived from the catalogue and presented in Section 4. We perform a detailed spectral index ($\\alpha^{5.5 GHz}_{1.4GHz}$) analysis, and identify ultra-steep and Gigahertz Peaked Spectrum sources in Section 5.  Concluding remarks are given in Section 6. We assume a Hubble constant of 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}$ = 0.27 and $\\Omega_{\\Lambda}$ = 0.73 throughout this paper. ",
        "conclusions": "We have presented new observations at 5.5 GHz of the extended Chandra Deep Field South. The 0.25 deg$^2$ region was observed with the Australia Telescope Compact Array using a mosaic of 42 pointings. The resultant image reaches an almost uniform noise level of $\\sim$12 $\\mu$Jy rms  and has a resolving beam of 4.9 $\\times$ 2.0 arcsec. Using a false-discovery-rate method, we extracted 123 sources or 142 source components. Ten sources were resolved multiple sources with jet-lobe structure and hence fitted as multiple components.  After carefully correcting for completeness, flux boosting and resolution bias, we derived source counts at 5.5 GHz. These are amongst the deepest source counts at 6cm but from an area 3 to 5 times larger than the previous counts to these depths. The ATLAS 5.5 GHz counts are consistent with the counts derived from other 5 GHz surveys at brighter flux densities, but are lower than counts in the literature by a factor of two for $S_{5.5GHz} < 0.3$ mJy. This discrepancy is attributed to cosmic variance because of the small size of the surveys involved. This fluctuation in the 5.5 GHz source counts at the faint end is similar to that seen at 1.4 GHz for $S_{1.4GHz} < 0.1$ mJy \\citep{norris2011}  The 1.4 -- 5.5 GHz spectral index has also been determined for all the sources in this field. We find a median spectral index for the ATCA 5.5 GHz sample of $\\alpha_{\\rm med} = -0.40$. We find the median spectral index for the faintest flux density levels in our sample, $S_{5.5 GHz} < 0.1$ mJy, is $\\alpha_{\\rm med} = -0.68$. The 1.4 GHz sources start to be missed in the 5.5 GHz image for $S_{1.4 GHz} < 0.5$ mJy so we performed a stacking analysis to examine the average spectral index of the full faint 1.4 GHz selected sample. From the stacking analysis we find that the faintest 1.4 GHz sources  ($40 < S_{1.4 GHz} < 200$ $\\mu$Jy)  also have steep spectra. This is consistent with the results of \\cite{owen2009} who found that radio spectra steepens at the faintest flux density levels ($S_{1.4 GHz} < 100$ $\\mu$Jy), after flattening from mJy to sub-mJy flux densities.  The 5.5 GHz sources have been cross-matched to existing optical/NIR data to obtain redshifts and we find, consistent with earlier work, that the $z-\\alpha$ relation seems to persist to these low flux densities. In addition, several candidate young AGN have been identified using the peak in the spectral energy distribution between 1.4 GHz and 6 GHz, suggesting that GPS sources are as common in the mJy population as they are at Jy levels.  In future work we will explore the spectral index properties of the ATLAS sources against other attributes, such as polarization and X-ray hardness, with the aim of distinguishing low-luminosity AGN from starforming galaxies in the sub-mJy population and determining how the observed properties of the sub-mJy population are linked to these two physical emission processes."
    },
    "1208/1208.0039.txt": {
        "abstract": "We compare the noise in interferometric measurements of the Vela pulsar from ground- and space-based antennas with theoretical predictions. The noise depends on both the flux density and the interferometric phase of the source. Because the Vela pulsar is bright and scintillating, these comparisons extend into both the low and high signal-to-noise regimes. Furthermore, our diversity of baselines explores the full range of variation in interferometric phase. We find excellent agreement between theoretical expectations and our estimates of noise among samples within the characteristic scintillation scales. Namely, the noise is drawn from an elliptical Gaussian distribution in the complex plane, centered on the signal. The major axis, aligned with the signal phase, varies quadratically with the signal, while the minor axis, at quadrature, varies with the same linear coefficients. For weak signal, the noise approaches a circular Gaussian distribution. Both the variance and covariance of the noise are also affected by artifacts of digitization and correlation. In particular, we show that gating introduces correlations between nearby spectral channels.   % % %We compare theoretical predictions of the scaling of noise with source intensity, %as a function of the flux density and interferometric phase of the source. %%We present measurements of the noise in interferometric observations of a  scintillating pulsar, %%the Vela pulsar, using ground- and space-based antennas. %The Vela pulsar provides a good laboratory for such studies because its %strong scintillation produces %intensity ranging from near zero to much greater than background noise, %and large phase variations on long baselines. %%We sample these spectra on time and frequency scales less than those of scintillation, %%in gates synchronous with the pulsar period. %We estimate noise from %comparison of samples closer together than the characteristic scales of scintillation. %In agreement with theoretical expectation, we find that %the noise is drawn from an elliptical Gaussian distribution in the complex plane, %centered on the signal. %The major axis lies at the phase of the signal, with the minor axis at quadrature. %The variance along the minor axis varies linearly with signal amplitude. %The variance along the major axis varies quadratically, %with the same linear terms as the minor axis. %At zero signal, noise follows a circular Gaussian distribution. %The variance of that distribution depends upon the signal level at digitization %and consequently upon whether or not other parts of the spectrum contain signal. %The noise depends upon the number of pulsar pulses within the sampling time. %We find that noise is slightly correlated between spectral channels, %in a pulse gate empty of signal, because of uneven sampling of the correlation function %due to pulsar gating. %% Kurtosis of noise in the absence of signal is small. ",
        "introduction": "Radioastronomical observations yield a deterministic part, the signal; and a random part, noise \\citep{tms86}. Consequently, the signal-to-noise ratio, the magnitude of the deterministic part divided by the standard deviation of the random part, characterizes them. An understanding of the noise is fundamentally important because it provides a measure of the possibility of detecting a weak signal, and of the reliability of the measurements of a detected signal. The noise is particularly important in situations where the signal varies, because the noise can mimic the signal. %In situations where the signal varies, the noise is important because it %can mimic the signal.  Noise includes background noise from the instrument and sky. However, because all radioastronomical signals are noiselike (with one possible exception: \\citet{jen01,Smi03}), they also contribute \\emph{self-noise} \\citep{kul89,ana91,viv91,mcc93,gwi06,gwi11a}. %All radioastronomical signals are noiselike %(with the exception of pulsar signals under certain circumstances: \\citet{jen01}). %The noiselike signal thus contributes to the noise in a measured spectrum. As argued previously \\citep{gwi06,gwi11a} and in Section\\ \\ref{sec:noise_dist_theory} below, for interferometric visibility, the variance of the noise increases quadratically with the signal in phase with the signal, and linearly with the signal at quadrature to the signal. The constant and linear coefficients in these two directions are equal.  We tested this picture for the cross-power spectrum of a scintillating source -- the Vela pulsar. These observations provided an ideal laboratory for such studies because the pulsar varies greatly in flux density and interferometric phase with frequency and time, because of interstellar scintillation \\citep{des92}. Thus, each spectrum spanned many scintillation elements. Furthermore, the signal-to-noise ratio of the strongest spectral peaks can be high, even for short integrations, and the source contributes significantly to system temperature, so that self-noise is important.  %In particular, the electric field of the source is drawn from a Gaussian distribution. %Through the action of the central limit theorem, %integration over frequency and time will drive the distribution of noise to a Gaussian distribution. %Such a distribution is fully characterized by variances and covariances. %Because of self-noise, %the variance depends on signal amplitude; with different forms in phase with the signal and at quadrature \\citep{gwi06,gwi11b}.  We compared predictions with observational estimates, formed by differencing samples close together in time and binning them according to their estimated average visibility. This procedure provides for convenient visualization of the noise distribution. Our investigation extends our previous results \\citep{gwi11b} to the regime of high signal-to-noise ratio and large interferometric phase variations. Because of pulsar gating, the number of samples per integration time was small and depended on the number of gated pulses within the integration time. From correlation functions, we determined the covariances of noise among spectral channels. Such covariances can be produced by quantization and by correlator effects.  We discuss these effects and compare results with observations.  \\subsection{Organization of this Paper}  Because the noise in the spectrum is drawn from a nearly Gaussian distribution with zero mean, variances and covariances characterize it. In Section\\ \\ref{sec:theory_bckgnd}, we introduce the theoretical basis for noise in interferometric visibility and present the mathematical descriptions used in the paper. In Section\\ \\ref{sec:observations}, we discuss our observations, correlation, and initial data processing. The remainder of the paper analyzes the noise in these observations. %We discuss our measurements of noise in the remainder of the paper. In Section\\ \\ref{sec:noise_dist_obs}, we quantify the distribution of noise and the influence of a signal. In Section\\ \\ref{sec:correlator_effects}, we discuss quantization and correlator effects. In Section \\ref{sec:conclusions}, we summarize our results and discuss implications for future observations and instruments. ",
        "conclusions": "\\label{sec:conclusions}  \\subsection{Summary of Results}\\label{sec:summary}  We compare theoretical predictions for the distribution of noise for cross-power spectra with observations of a scintillating pulsar, the Vela pulsar. We describe observations made with Earth-based VLBI baselines, and with baselines from an orbiting spacecraft to Earth. These observations extend previous studies \\citep{gwi11a} to the regime of high signal-to-noise ratio and large variations in interferometric phase.  In Section\\ \\ref{sec:theory_bckgnd}, we argue that, in the presence of signal, noise on a short baseline should be drawn from an elliptical Gaussian distribution in the complex plane. The theory was previously presented in \\citet{gwi06} and \\citet{gwi11a}. The major axis of the distribution is aligned with the direction of the signal. The variance along the minor axis depends linearly on signal strength; the variance along the major axis has the same linear dependence, plus a quadratic term. At zero signal, the major and minor axes are equal and the distribution of noise is a circular Gaussian, as for a gate or spectral region empty of signal.  We test this theory with observations on the baselines from Mopra, Hartebeesthoek, and the VSOP spacecraft to Tidbinbilla in Section\\ \\ref{sec:noise_dist_obs}. We estimate noise by comparing samples within the characteristic scales of the scintillation, and binning their differences by average interferometric visibility. We find that the distribution of noise closely follows the expected elliptical Gaussian form for each visibility, and the scaling with visibility of the major and minor axes corresponds to the quadratic and linear noise polynomials, respectively. The quadratic coefficient accurately reflects the number of samples and the contribution of intrinsic amplitude variations; the constant coefficient agrees with that estimated from empty portions of the spectrum for the Mopra-Tidbinbilla and Hartebeesthoek-Tidbinbilla baselines.  In Section\\ \\ref{sec::Correlator_Effects}, we demonstrate that quantization, gating, and integration each affect the properties of the noise. One interesting consequence is that the noise in the presence of signal is less than that in a completely empty spectrum -- a result of the combination of quantization and spectral variations. Also, pulsar gating leaves fewer samples, thus larger noise, at larger lags; this effect incurs correlations in the spectral noise. In principle, complete knowledge of the quantizer levels for each integration period, and of the autocorrelation functions at the two antennas, allows calibration of these effects.% \\citep[see][]{jen98}. Alternatively, recording the signal with many quantizer levels increases the dynamic range, and reduces effects of variation in quantizer levels. Flexible software correlators, such as the DiFX correlator \\citep{del07}, can control these artifacts, while Nyquist-sampled spectra of individual pulses obviate the difficulties in characterizing inhomogeneities within the integration \\citep{Joh12a}.   \\subsection{Self-Noise for an Interferometer}  We present observations for a scintillating pulsar, but the effects of self-noise hold for any interferometric observation. A careful evaluation of these effects is essential for a priori estimates of the accuracy of pulse-timing and spectroscopy using single-dish observations, and for scintillation studies and astrometry using interferometry.  Many telescopes now under construction or being planned, such as LOFAR, ASKAP, SKA, and so on will operate as interferometers, with many baselines among many antennas of a particular design. Each baseline will have the distribution of noise we describe above in Section\\ \\ref{sec:theory_bckgnd}, and as we observe for the Vela pulsar. Because each telescope in a large array receives precisely the same noiselike signal from the source, increases in the number of antennas and averaging of many baselines do not change self-noise, when expressed in terms of flux density. However, averages over many baselines do decrease the background noise. Stated in terms of the notation introduced in Section\\ \\ref{sec:theory_bckgnd}, increasing the number of identical antennas $N_A$ reduces $b_0$  as $N_A^{-2}$ and $b_1$  as $N_A^{-1}$ but does not change $b_2$. For true ``tied-array'' operation, where electric fields from all antennas are phased and summed before correlation, statistics are those of a single dish, \\citep[][Equation\\ 11]{gwi11a}. As the number of antennas become larger, self-noise becomes more important. When the source dominates the system temperature, further improvements demand more samples $N_{\\rm obs}$, as produced by wider bandwidth or longer integration time: a greater aperture does not provide more accuracy. Our expression Equation\\ \\ref{eq:noise_perp_para} generalizes this result to interferometry.  Astrometry depends on measurements of interferometric phase. Equations\\ \\ref{eq:noise_perp_para} and \\ref{eq:dicke_complex} show that the maximum attainable phase accuracy is $\\delta\\phi \\approx \\sigma_{\\perp}/s \\approx \\sqrt{n/(N_{\\rm obs} s)}$, or the inverse square root of the signal-to-noise ratio $s/n$, divided by $\\sqrt{N_{\\rm obs}}$. Similarly, the maximum attainable accuracy in measurement of flux density by an interferometer, or a large single dish, is approximately the flux density of the source divided by $\\sqrt{N_{\\rm obs}}$. Likewise, the maximum attainable accuracy in pulsar timing is approximately the width of the narrowest feature in the profile, divided by the signal-to-noise ratio and by $\\sqrt{N_{\\rm obs}}$.  However, when self-noise is the limiting factor, the maximum attainable accuracy is simply the width of that feature, divided by $\\sqrt{N_{\\rm obs}}$."
    },
    "1208/1208.1411_arXiv.txt": {
        "abstract": "We explore observables in a lattice Universe described by a recently found solution to Einstein field equations. This solution models a regular lattice of evenly distributed objects of equal masses. This inhomogeneous solution is perturbative, and, up to second order in a small parameter, it expands at a rate exactly equal to the one expected in a dust dominated Friedmann-Lema\\^itre-Robertson-Walker (FLRW) model with the equivalent, smoothed, energy density. Therefore, the kinematics of both cosmologies are identical up to the order of perturbation studied. Looking at the behaviour of the redshift and angular distance, we find a condition on the compactness of the objects at the centre of each cell under which corrections to the FLRW observables remain small, i.e. of order of a few percents at most. Nevertheless, we show that, if this condition is violated, i.e. if the objects are too compact, our perturbative scheme breaks down as far as the calculations of observables are concerned, even though the kinematics of the lattice remains identical to its FLRW counter-part (at the perturbative order considered). This may be an indication of an actual fitting problem, i.e. a situation in which the FLRW model obtained from lightcone observables does not correspond to the FLRW model obtained by smoothing the spatial distribution of matter. Fully non-perturbative treatments of the observables will be necessary to answer that question. ",
        "introduction": "It has long been recognised that calculating observables in an inhomogeneous Universe could be quite challenging. Indeed, our usual description of the geometry of the Universe on large scales relies on Friedmann-Lema\\^itre-Robertson-Walker (FLRW) models in which the distribution of matter is assumed perfectly homogeneous and isotropic. As a result, light rays propagate in a homogeneous medium filled with matter and are therefore sensitive only to the Ricci curvature of spacetime. In the real Universe, on the other hand, at least in the late stages of its evolution, matter is clumped into virialised objects with large (almost) empty regions between them, and light therefore travels mainly in empty space, where its behaviour is dominated by the Weyl curvature of spacetime, rather than the Ricci curvature. This raises the natural question to whether the FLRW approximation is suitable to calculate observables in the late time Universe: under which conditions is it possible to replace Weyl curvature along the line of sight with an 'equivalent' Ricci curvature? And what is this 'equivalence' all about? The problem is therefore what has been dubbed a 'fitting problem' \\cite{Ell84, Ellis:1987zz}, that is: how do we replace the real Universe by an FLRW idealisation? In principle, this leads to a backreaction issue: the idealised FLRW given by observations will not, in general be the FLRW model for which the actual matter density has simply been smoothed in space. In the standard model, this is usually accounted for via the Dyer-Roeder equation \\cite{Dyer:1973zz} (see \\cite{1969ApJ...155...89K} for earlier results on this problem), but recent works tend to show that there might be better ways to model the effect \\cite{Clarkson:2011br, Bolejko:2010nh,Kainulainen:2009dw,Kainulainen:2010at}. This issue has been addressed by a certain number of authors lately; see e.g., \\cite{Bolejko:2010nh, Meures:2011gp,Clarkson:2011br,Bolejko:2012ue}. Usually, they use a 'realistic' model of structure formation, either through cosmological perturbation theory or N-body simulation, to evaluate the impact of inhomogeneities on the propagation of light; others have used Swiss-Cheese models \\cite{Brouzakis:2007zi,Marra:2007pm,Vanderveld:2008vi,Flanagan:2011tr}, mainly to see whether inhomogeneities along the line of sight could account for the Dark Energy phenomenon. Our take on the problem is slightly different. In particular, we are not trying to address the Dark Energy problem; we simply try to understand under which conditions, in the controlled environment of an (almost) exact solution to Einstein field equations, the behaviour of null geodesics can safely be approximated by the null geodesics of an FLRW model.\\\\ \\\\ In a previous paper \\cite{Bruneton:2012cg}, we proposed a lattice solution to Einstein field equations made of equal masses $M$ separated by a comoving distance $L$. This solution is accurate at order $M/L$, and can be expanding or contracting. We proved that this solution is kinematically equivalent (at the order of perturbation considered) to an FLRW model with a dust matter content having an energy density equal to the one obtained by smoothing the lattice distribution, i.e. $M/L^{3}$. This was interpreted as supporting the usual fluid approximation in cosmology. Recently cosmological lattices similar to the one presented in \\cite{Bruneton:2012cg} have been studied both analytically \\cite{Clifton:2012qh} and numerically \\cite{Yoo:2012jz}. The lattice studied in \\cite{Yoo:2012jz} is similar to the one we studied, with cubic symmetry, and their results agree with ours, since they recover an Einstein-de Sitter Universe in the limit $M/L\\ll 1$ (large separation between the masses). Similarly, the results of \\cite{Clifton:2012qh} show that, for an infinite number of masses, the kinematics of the lattice tends towards the kinematics of the Einstein-de Sitter Universe.\\\\ The solution in \\cite{Bruneton:2012cg}, being kinematically equivalent to an FLRW model at the perturbative order considered, is the ideal setting to study observables: if observables only deviate slightly from the observables in the kinematically equivalent FLRW Universe, then we might say that this FLRW model is a good fitting model. On the other hand, if observables exhibit large variations compared to the one in the analogue FLRW, that will show that a perturbative calculation is unable to describe accurately observables in the lattice, even though it is a good approximation to the geometry of spacetime; in that case, further, non perturbative studies will be needed.\\\\ \\\\ The propagation of light in lattice models has been studied before \\cite{Clifton:2009jw,Clifton:2011mt} in the context of Lindquist-Wheeler models \\cite{LindWheel}. These models are only approximate, and the propagation of light through the boundaries between cells is not fully controlled, leading to important differences in the results, depending on the approximation scheme used \\cite{Clifton:2011mt}. Our model, on the contrary, is only approximate in power of $M/L$, and not on the way cells are glued together (much like the model developed in \\cite{Clifton:2010fr}). Therefore, propagating light through the lattice is not a problem. The only limitation will come from the fact that we can only trust the solution for a small range of redshifts (typically $z\\ll 1$). By solving Sachs equations at order $M/L$, we will see that at this order the equations for the shear and the isotropic expansion decouple. This implies that at order $M/L$, {\\it a priori}, the Weyl curvature of spacetime does not play any role in the calculation of the distance/redshift relation. Then, the formal solution appears to be equivalent to the solution in the analogue FLRW model, plus small corrections of order $M/L$. Nevertheless, by studying carefully the order of magnitude of these corrections, we will prove that they remain small only under specific conditions on the compactness of the objects forming the lattice: when these objects are too compact, the perturbative expansion breaks down, and the differences between observables in the lattice and in the analogue FLRW model cannot be simply evaluated by using perturbative methods. That might indicate a tension between the fitting models constructed kinematically and observationally. This paper somehow illustrates, on a specific example, the concepts defined in \\cite{Kolb:2009rp}: using the terminology and classification introduced in that paper, in our lattice Universe, the strong backreaction is absent and the Global and Averaged Background Solutions (GBS and ABS) coincide at the perturbative order considered here, but weak backreaction are present and the Phenomenological Background Solution (PBS) does not necessarily coincide with the others. Indeed, on the one hand, if our compactness criterion is satisfied, the difference between the PBS and the GBS (and ABS) is small, of order of a few percents at most. On the other hand, when objects in the lattice are too compact, perturbative calculations cannot be trusted to give the structure of the PBS, and fully non-perturbative calculations must be done to evaluate by how much this PBS differs from the ABS (and from the GBS). The break-down in the perturbative calculations can be attributed to the fact that the shear of the ray bundles cannot be neglected anymore.\\\\ \\\\ The paper is organised as follows. Section~2 will briefly present the solution described in details in \\cite{Bruneton:2012cg}. Section~3 details the calculations of the redshift and angular distance in the lattice Universe. For the sake of clarity, Section~3 focusses only on the analytical expressions, while their physical interpretation, numerical calculations, and discussion of the effect of inhomogeneities on the propagation of light, especially in regards of the fitting problem and of the Weyl focussing, are left for Section~4. Finally, section~5 will be a conclusion. \\\\%Some technical results will be left for an appendix.\\\\ \\\\ Throughout the paper, the signature of the spacetime metric will be $(-,+,+,+)$, and, unless otherwise specified, we will work in units of $c=G=1$. ",
        "conclusions": "In this paper, we showed that in a lattice Universe kinematically equivalent to an FLRW model with the same averaged energy density, thus showing no 'strong backreaction' in the sense of \\cite{Kolb:2009rp}, angular and luminosity distances would not deviate significantly from the ones in the FLRW model provided the spread $\\eta/L$ of the object was related to the mass of the objects $M/L$ via: \\beq \\label{mainboundConclusion} \\frac{M}{L} \\ll \\mathcal{O}(1) \\times \\left(\\frac{\\eta}{L}\\right)^4. \\eeq In our model, compactness is thus a key parameter in regards of the amplitude of the 'weak backreaction' \\cite{Kolb:2009rp}. In other words, this relation can thus be understood as a criterion to decide whether or not the fitting problem is a problem at all within these models: if this relation is satisfied, observables are almost the ones of the kinematically fitting FLRW model. Otherwise, if objects are too compact, perturbative estimates are not reliable, and it is impossible to say how observables relate to the fitting model without a fully non-perturbative treatment of the propagation of light. This should not be understood as a claim that there is a genuine fitting problem in our 'real' Universe, but simply as a warning that there exist spacetime configurations such that, despite the fact that the solution of Einstein field equations remains close to an FLRW configuration, observables might significantly deviate from their FLRW analogue, or at least, need to be calculated non-perturbatively; see \\cite{Clarkson:2011uk} for a similar result in FLRW plus perturbations. Indeed, we hinted at a link between this break-down of the perturbative expansion and the fact that the Weyl curvature behaves non perturbatively when the bound (\\ref{mainboundConclusion}) is not satisfied, leading to big corrections to the angular distance at second order in $\\sqrt{M/L}$ that might be compensated non perturbatively by the effect of the shear. It would be interesting to check if similar conclusions can be drawn in more 'realistic' configurations of the matter distribution. In particular, provided one could find a satisfactory way of estimating the compactness of cosmological objects, it would be interesting to see whether or not the bound (\\ref{mainboundConclusion}) is satisfied in our Universe, using galaxy surveys and/or N-body simulations. A future work must also present the numerical solution of the full system of Sachs equations non perturbatively, in order to avoid the limitations of the bound (\\ref{mainboundConclusion}); only such a solution will allow one to decide what happens when this bound is not satisfied. The results of \\cite{Holz:1997ic} imply, through statistical arguments, that there is no fitting problem in cosmology, even for a matter distribution made of point masses, except for the exceptional light-rays that travel too close to some masses. Our model, with a high degree of symmetry, seems to indicate otherwise: the 'quasi-poles' seem to act only in one way, decreasing the angular distance compared to its FLRW counter-part, and their effect cannot be accounted for by a simple encounter of the light ray with the neighbourhood of a mass. Nevertheless, it is actually impossible to conclude and one will need the full, non perturbative, solution in order to do some statistics and to tests the results of \\cite{Holz:1997ic} with the lattice solution presented here. This is the subject of an ongoing investigation.  \\ack J.-P. B. is FSR/COFUND postdoctoral researcher at naXys. J.-P. B. and J. L. acknowledge fruitful discussions with T. Clifton while preparing this paper. We thank anonymous referees for their valuable comments on an earlier version of this paper."
    },
    "1208/1208.1319_arXiv.txt": {
        "abstract": "We consider the bending of light by nonlinear electrodynamics when the magnetic field $B$ exceeds the critical value $B_{\\rm c} = m^2 c^2 /e \\hbar = 4.4 \\times 10^9 \\rm T$. Using the index of refraction derived from the analytic series representation in one-loop effective action of QED, we found the trajectory and the bending angle of light in geometric optics. The angle bent by ultra-strong magnetic field of magnetar was estimated and compared with the gravitational bending. The result may be useful in studying the lensing, birefringence, and other nonlinear quantum electrodynamic effects above  $B_{\\rm c}$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.4199_arXiv.txt": {
        "abstract": "{We present collapse simulations of strongly magnetised, 100 M$_{\\odot}$, turbulent cloud cores. Around the protostars formed during the collapse Keplerian discs with typical sizes of up to 100 AU build up in contrast to previous simulations neglecting turbulence. Analysing the condensations in which the discs form, we show that the magnetic flux loss is not sufficient to explain the build-up of Keplerian discs. The average magnetic field is strongly inclined to the disc which might reduce the magnetic braking efficiency. However, the main reason for the reduced magnetic braking efficiency is the highly disordered magnetic field in the surroundings of the discs. Furthermore, due to the lack of a coherently rotating structure in the turbulent environment of the disc no toroidal magnetic field necessary for angular momentum extraction can build up. Simultaneously the angular momentum inflow remains high due to local shear flows created by the turbulent motions. We suggest that the \"magnetic braking catastrophe\" is an artefact of the idealised non-turbulent initial conditions and that turbulence provides a natural mechanism to circumvent this problem.} ",
        "introduction": "\\label{sec:intro}  In recent years a number of authors have studied the formation of protostellar discs during the collapse of strongly magnetised magnetic cloud cores~\\cite{Hennebelle08a,Seifried11}. In simulations with magnetic field strengths comparable to observations no rotationally supported discs were found. As strong magnetic braking is responsible for the removal of the angular momentum, this problem is also called the \"magnetic braking catastrophe\". The results of these simulations stand in contrast to observations which show that discs are present in the earliest stage of protostellar evolution. Here we present results from a number of simulations investigating the role of turbulence in reducing the magnetic braking efficiency and allowing for the formation of protostellar discs.  We now shortly describe the basic simulation setup (see Seifried et al. 2012 for a more detailed description). We simulate the collapse of a 100 M$_\\odot$ molecular cloud core which is 0.25 pc in size, threaded by a magnetic field in the z-direction -- the mass-to-flux ratio is $\\mu = 2.6$ -- and rotating around the z-axis. Additionally, we add a supersonic turbulence field with a power-law exponent of p = 5/3. The turbulent energy is equal to the rotational energy, i.e. $\\beta_{turb}$ = 0.04, corresponding to a turbulent rms-Mach number of $\\sim$ 2.5. We performed several simulations with different turbulence seeds to test the dependence of our results on the initial conditions. We find that the results do not depend on the chosen initial turbulence field. ",
        "conclusions": "We have observed the formation of rotationally supported discs during the collapse of massive, magnetised and turbulent cloud cores. We attribute this to the turbulent surroundings of the discs which reveal a highly disordered magnetic field and no signs of a coherent rotation structure. In contrast, magnetic flux loss is not able to account for the formation of Keplerian discs as only on the disc scale ($\\le$ 100 AU) significant flux loss is observed. However, we emphasise that the magnetic braking efficiency has to be reduced already on larger scales ($\\ge$ 500 AU) even before the gas falls onto the disc. This is demonstrated by a corresponding non-turbulent run where the angular momentum is removed largely already outside the disc. Furthermore, non-ideal MHD effects, which have already been shown that they cannot account for the formation of early-type Keplerian discs, are not required as turbulence alone provides a natural and simple mechanism to circumvent the ``magnetic braking catastrophe''. Hence, our work strongly suggests that the magnetic braking problem as reported in numerous papers is more or less a consequence of the highly idealised initial conditions neglecting turbulent motions."
    },
    "1208/1208.1775_arXiv.txt": {
        "abstract": "The orbital angular momentum of a close-orbiting giant planet can be sufficiently large that, if transferred to the envelope of the host star during the red giant branch (RGB) evolution, it can spin-up the star's rotation to unusually large speeds. This spin-up mechanism is one possible explanation for the rapid rotators detected among the population of generally slow-rotating red giant stars. These rapid rotators thus comprise a unique stellar sample suitable for searching for signatures of planet accretion in the form of unusual stellar abundances due to the dissemination of the accreted planet in the stellar envelope.  In this study, we look for signatures of replenishment in the Li abundances and (to a lesser extent) \\cratio, which are both normally lowered  during RGB evolution. Accurate abundances were measured from high signal-to-noise echelle spectra for samples of both slow and rapid rotator red giant stars. We  find that the rapid rotators are on average enriched in lithium compared to the slow rotators, but both groups of stars have identical  distributions of \\cratio\\ within our measurement precision. Both of these abundance  results are consistent with the accretion of planets of only a few Jupiter masses. We also explore alternative scenarios for understanding the most Li-rich stars in our sample---particularly Li regeneration during various stages of stellar evolution. Finally, we  find that our stellar samples show non-standard abundances even at early RGB stages, suggesting that initial protostellar Li abundances and \\cratio\\ may be more variable than originally thought. ",
        "introduction": "\\label{sec:intro} Composition studies of red giant stars have yielded insightful clues into the internal structure and history of low mass stars.  Some elemental abundances do not change much over the entire lifetime of the star, and these abundances can encode information about the star's birthplace and the chemical evolution history of the gas that became the star. Other abundances, particularly those of light elements, are altered by nuclear processes throughout the star's life, and measurements of these abundances and isotopic ratios have  provided observational constraints on stellar evolution models. Light abundance alterations also provide a means of studying planet accretion in red giant stars. The abundances of some light element in red giants are significantly depleted from the stars' initial stellar abundances and, by extension, from the abundances of planets that formed from the same protostellar material.  Therefore, accreted planets can replenish the stellar surface abundances of depleted elements. In this work, we focus on the lithium abundance,  \\lli \\footnote{The abundance of Li is defined as $A({\\rm Li})\\equiv\\log N_{\\rm Li}-\\log N_{\\rm H} +12.00$}, and the ratio of $^{12}$C to $^{13}$C. The initial stellar and planetary abundances of these elements in solar metallicity stars, as measured in the solar system, are \\lli$\\sim3.3$\\, dex \\citep[see, e.g., ][]{bosegaard88,lodders98} and \\cratio$\\sim$90 \\citep{anders89}.  Lithium is  destroyed at relatively low temperatures  in a star by proton capture reactions, and its abundance both in the stellar interior and at the stellar surface changes considerably throughout the lifetime of most stars. During the main sequence (MS) phase, Li is preserved only in the coolest, outermost regions of the star that comprises only a few percent of the total stellar mass.  In the stellar interior, the star is essentially devoid of Li. For stars more massive than $\\sim1.4$\\,\\msun, the surface \\lli\\ stays constant throughout  the MS lifetime, while \\lli\\ slowly decreases in lower mass stars. Therefore, by  the end of the MS, the surface \\lli\\  in the outer convection zone can range from near the original stellar abundance to depleted by factors of 100, depending on the star's mass. Near the beginning of the post-MS evolutionary phases, the star undergoes first dredge-up (FDU), during which the convection zone first deepens into the star and then recedes. The Li that had been preserved  in the outer $\\sim1$\\% of the stellar mass on the MS gets mixed into the Li-poor interior, diluting the surface abundance by a factor of $\\sim 60$ \\citep{iben67}. Similar processes affect the \\cratio\\ ratio. The main chain of the CNO cycle reduces \\cratio\\ in the stellar interior towards the equilibrium value of 3.5, and mixing during FDU will lower the overall  \\cratio\\ of a red giant's envelope by a factor of $\\sim 3$ \\citep{sweigart89}. These are the ``standard model'' abundance changes, which have been observationally verified by, e.g.,  \\cite{lambert81} and \\cite{brown89}.  Stars less massive than  $\\sim$2.3\\,\\msun\\ will experience additional surface abundance changes when the outward moving H-burning shell reaches the chemical discontinuity left behind by the convective envelope during FDU, and this evolutionary stage creates in the color-magnitude diagram what is known as the luminosity bump or RGB bump. In stars more massive than $\\sim$2.3\\,\\msun, the H-burning shell does not reach the chemical discontinuity before the star evolves off of the red giant branch (RGB). The ``erasing'' of the chemical discontinuity in low mass stars allows for additional  non-convective mixing methods to further decrease the surface \\lli\\ and \\cratio\\ by varying amounts. \\cite{denissenkov03} coined the term ``canonical extra-mixing'' to refer to the prevalent processing of light elements beyond standard dilution predictions for low-mass, post-RGB bump giant stars that have been observed  in open clusters \\citep{gilroy89,gilroy91,luck94,smiljanic09}, globular clusters \\citep{shetrone03,pilachowski03,reicoblanco07}, and field red giants \\citep{sneden86,charbonnel98,gratton00,kumar09}. The study by \\cite{denissenkov03} described  a two-parameter turbulent diffusion model \\citep[following the work by ][]{wasserburg95,boothroyd99} and explored rotation-driven mixing  \\citep[as previously employed by ][]{sweigart79,charbonnel98} as a physical source of the turbulent diffusion. They concluded that most of the non-standard abundances seen in red giants could be reproduced by a relatively simple two-parameter diffusion model.   However, within this framework of standard dilution and canonical extra-mixing models, inexplicable light-element  abundances can remain. Of particular interest to the present study are the  few percent of red giants that are ``Li-rich,'' with  \\lli\\,$>1.5$\\,dex \\citep{dasilva95,balach00,drake02,reddy05,kumar09,gonzalez09Li,carlberg10b,kumar11,monaco11,ruchti11,lebzelter12}.  The Li abundances of these stars are larger than that expected from stellar evolution and sometimes even exceed the assumed undiluted abundances of these stars \\citep[e.g.,][]{balach00,kumar09,carlberg10b,monaco11,ruchti11}. For massive, luminous red giant stars (i.e., $M_{\\star}>1.5$\\,\\msun and $\\log L/L_{\\sun}>4$), temperatures at the base of the convection envelopes are hot enough for the nucleosynthesis of $^7$Li through the Cameron--Fowler chain \\citep{cameron71}.   If this freshly synthesized Li (or the $^7$Be that decays to Li) is quickly transported to the cooler  regions of the convective envelope, then the surface \\lli\\ can substantially increase \\citep{scalo75}. This mechanism is known as hot bottom burning (HBB). However, many of the Li-rich giants do not fit this picture because their convective envelopes are too cool, even at the hottest depths, for Li synthesis. Both \\cite{charbonnel00} and \\cite{reddy05} found that many of these inexplicably Li-rich giants tend to cluster near the luminosity bump.   It is thought that the erasing of the chemical discontinuity at the bump may allow non-convective mixing processes to connect the cool convective envelope to  hotter depths, where Li can be regenerated. Like the HBB mechanism, the mixing must be rapid to replenish  the surface abundance. For example, \\cite{sack99} used a parameterized ``conveyor belt'' model and could reproduce the abundances of Li-rich giants with certain mixing geometries; however, they did not provide a physical mechanism for their successful models. Rotation was thought to be a likely mechanism \\citep{denissenkov04} until \\cite{palacios06} found that a self-consistent model of rotational mixing could not generate enough circulation to build up Li in the stellar envelope. Recently, \\cite{palmerini10} summarized various mixing mechanisms capable of providing the needed conveyor belt and noted that the two current contenders are thermohaline mixing \\citep{eggleton06,charbonnel07} and magnetic buoyancy \\citep{busso07,nordhaus08,guandalini09,palmerini10}. Because thermohaline mixing is a relatively slow process, it may only be an alternative model for the canonical extra-mixing that destroys Li. The magnetic buoyancy models circulate material fast enough to replenish Li in the stellar envelope; however, both the  \\cite{guandalini09} and \\cite{palmerini10} models predict maximum lithium enrichments of \\lli$\\sim$\\,2.5 dex, which  are an order of magnitude smaller than the abundances observed in some of the most Li-rich red giants.   Although there is no consensus on the physical mixing mechanisms capable of transporting freshly regenerated Li to the stellar envelope, the intrinsically nuclear origin of the Li can be observationally distinguished from Li supplied by an accreted planet. First, any extra mixing should dredge-up  material with low \\cratio, which will further decrease the surface \\cratio.  The ``enhanced extra-mixing\" models of \\cite{denissenkov04} yield \\cratio\\ between 8.2 and 17.4 compared to the expected values of 23 without the enhanced mixing.  Planet accretion, by contrast, should raise both \\lli\\ and \\cratio. Second, because only a small fraction of stars associated with the RGB bump are Li-rich, then either only a small fraction of stars evolving through the bump regenerate Li, which is then long-lived in the envelope, or all stars evolving through the bump phase experience a short-lived phase of high Li abundances---a ``Li flash.'' \\cite{charbonnel00} considered the latter scenario more likely because of the paucity of Li-rich giant stars between the luminosity bump and the RGB tip. Therefore, Li-rich giants at other evolutionary stages (particularly pre-bump stages) require an alternative explanation, such as planet accretion. Finally, accreted planets should contribute angular momentum to the stellar envelope in addition to altering the stellar abundances of light elements. \\cite{carlberg09} showed that some known exoplanets are expected to experience tidal orbital decay, and the orbital angular momentum that is transferred to the stellar envelope is sufficiently large to measurably increase the host stars'  rotational velocities.     Because red giant stars are generally slow rotators \\citep{deMed96b,gray81,gray82}, red giants with enhanced rotation and replenished light-element abundances are candidates for stars that have accreted a planetary companion.   In our project, we select  samples of both rapid rotators and slow rotators (as a control sample) to look for multiple signatures of planet accretion simultaneously. In addition to a Li enhancement, we are  looking for evidence of enhancements of \\cratio\\  and (in a future paper) refractory elements in the atmospheres of rapid rotators, which should also be indicative of planet accretion. Interpretation of these signatures will likely be open for debate  for any individual star; however, trends that differ significantly between the two main samples (defined only by their relative \\vsini) should make a much stronger case to either defend or refute the proposition that some rapid rotators gain their angular momentum from a former planet. Ours is not the first attempt to correlate Li-rich giants and rapid rotation nor the first to invoke planet accretion as the explanation for the correlation. Planet accretion was first put forward by \\cite{alexander67} to explain Li enrichment in giant stars, while \\cite{pete83} were the first to consider planets as sources of angular momentum in evolved stars.  \\cite{drake02} noted that while both slow and rapid rotators could be Li-rich, only 2\\% of the slow rotators were Li-rich compared to nearly 50\\% of the rapid rotators.  \\cite{siess99} modeled the  accretion of sub-stellar companions by red giant stars and calculated observational signatures of planet  accretion in those stars, such as rapid rotation and $^7$Li enrichment. From the actual occurrence of these predicted observational signatures in the red giant population, \\cite{siess99} estimated that  4\\%--8\\% of Sun-like stars host giant planets at orbital radii small enough for significant interactions to take place between the stars and  their planets while  the stars are on the RGB. Others who have considered  planet accretion as the solution to both Li enhancement and rapid rotation in giants include \\cite{wallerstein82}, \\cite{reddy02a}, \\cite{carney03}, and \\cite{denissenkov04}.  We describe the selection and observation of our red giant sample in Section \\ref{sec:sample_obs}. In Section \\ref{sec:stellparams}, we detail the derivation of the atmospheric properties, rotational velocity, and abundances. We present our combined abundance and rotation results in Section \\ref{sec:placc}, where we  find  Li enhancement in the rapid rotators that is consistent with planet accretion. We find an unmeasurable difference in \\cratio\\ between the slow and rapid rotators; however, this \\cratio\\ result is still consistent with planet accretion given our measurement uncertainties and the modest \\cratio\\ enhancement expected to accompany the observed level of Li enrichment. We divide our sample into groups of similar evolutionary stages and find that the Li enrichment of the  rapid rotators persists within these subgroups. In Section \\ref{sec:discuss}, we discuss alternative internal and external mechanisms that could explain enriched Li and how the enhanced angular momentum of the rapid rotators might affect these mechanisms. We select a small group of Li rich rapid rotators that are the best candidates for planet accretion to test whether the rotation and Li enrichment is consistent with planet engulfment. Our conclusions are presented in Section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discuss} In the previous section, we found two important  results. First, the rapid rotators tend to be more Li-rich than the slow rotators regardless of evolutionary stage. Second,  there are at least two stars with \\lli\\ exceeding even the most conservative estimates of Li dilution (the standard dilution models) as well as a handful of stars that appear to be Li-rich when comparing their abundances to the other stars of similar evolutionary stage, but whose abundances compared to the models in Figure \\ref{ali_cratio_new}  are not atypical. Here, we consider various scenarios for understanding these two results.  \\subsection{Post-bump Lithium Regeneration} The prevailing explanation for Li-rich stars is  internal Li regeneration via the Cameron-Fowler chain, which is thought to begin  at the luminosity bump when the convective envelope can be connected to the H-burning shell via non-convective (generally diffusive) mixing processes.   At the luminosity bump, the hydrogen burning shell erases the chemical discontinuity, or ``$\\mu$-barrier,'' that inhibits mixing. Considerable work has been put into trying to understand the physical causes of this ``cool-bottom processing'' (CBP) mechanism---whether the mixing  results from, e.g.,  a thermohaline instability or buoyant magnetic flux tubes. Regardless of the underlying physics, one observational  consequence of this explanation is that the star must have evolved to the luminosity bump, and indeed, \\cite{charbonnel00} found  most Li-rich giants to be associated with this evolutionary stage. Another observational  signature of this processing is the further reduction of \\cratio\\ that occurs after the Li regeneration ceases.  Stars with the largest \\lli\\ may appear to have normal \\cratio, while stars with more modest Li enrichments will show abnormally low \\cratio.   On average, therefore, \\cratio\\ should be lower in stars that have undergone Li regeneration. Considering these two observational signatures of Li regeneration, we can conclude that our Li-rich rapid rotators do not fit this picture for the following reasons. First, on average the rapid rotators do not exhibit lower \\cratio\\ than the slow rotators. Second, if our stellar evolution groups are defined correctly, then the two most Li-rich stars are at pre-bump stages and should not have  yet experienced  Li regeneration.  \\subsection{Rotation-induced Mixing} An alternative Li-regeneration solution is one in which the enhanced rotation of the rapid rotators drives the non-convective mixing that connects the envelope to the Li-burning interior. The assumption is that the rapid rotation creates favorable mixing currents that can deposit the newly regenerated Li (or the parent isotope $^{7}$Be) into the stellar envelope.  In this scenario, the star's evolutionary proximity to the luminosity bump is irrelevant because the enhanced rotation is itself responsible for the non-canonical mixing. Rotation-induced mixing has been explored by both \\cite{sack99} and \\cite{denissenkov04} to explain Li-rich giants, and both models are also capable of lowering \\cratio\\ from the standard model predictions. However, both groups were also exploring rotation-induced mixing as a possible CBP mechanism. In other words, their models required the erasure of the chemical discontinuity at the luminosity bump, whereas a successful model for explaining our pre-bump Li-rich rapid rotators is one which does not require the erasure of the mean molecular weight discontinuity to function. \\cite{chaname05} were the first to relax the assumption that the $\\mu$-barrier inhibits all extra mixing  to see whether rotational mixing would occur at other evolutionary phases. They explored the evolution of  different rotational profiles and found that rotational mixing does indeed become active beyond the luminosity bump but only if the CZ was  differentially rotating. In other words, in the case of solid body rotation in the stellar envelope, a star {\\it would not} experience extra mixing even after the $\\mu$-barrier was erased. Additionally, they found that the FDU mixing of \\cratio\\ occurred earlier and more gradually in their rotating models compared to standard models (though both models resulted in the same post-FDU value). On the other hand, to reproduce the abundance observations, the initial rotation rates  of their models must exceed the values actually observed. The \\cite{chaname05} paper did not explore the explicit effect of rotation on the Li abundance, in particular whether their models allow the special mixing cases capable of creating at least short-lived periods of enhance Li (as opposed to further destruction of Li that generally occurs with extra mixing).  Further work is needed to see whether such conditions are possible and whether those conditions require the removal of the ``$\\mu$-barrier'' or not.   \\subsection{Helium Flash} Recently,  \\citet[][hereafter, K11]{kumar11} conducted a large survey for Li-rich stars in  a sample of 2,000 red giants  and found 15 previously unidentified Li-rich stars. Like our study, they find  some  Li-rich giants that are too warm to be associated with the luminosity bump.  K11 suggest that these stars are  likely associated with the red clump---the core helium-burning stage of relatively metal-rich stars.  They suggest two scenarios for explaining the presence of large \\lli\\ in post-RGB stars. The first explanation  is simply the survival of Li that was regenerated at the luminosity bump; however, this scenario is  contradicted by the relatively few Li-rich stars between the bump and RGB tip.  K11 also hypothesize the possibility that Li could be regenerated  during the He flash, although this suggestion is based simply on the observed coincidence  of their  Li-rich giants near the red clump more so than on any physical mechanism. The idea is that some fraction of the $^3$He remaining in the  RGB tip star could be converted to Li through the Cameron-Fowler process, whereby $^3$He($^4$H,$\\gamma)^7$Be$(e^-,\\nu)^7$Li. They suggest that this model would only operate in stars within the narrow mass range of $1.5M_\\sun\\leq M_\\star \\leq 2.25 M_\\sun$,  where the star has a low enough mass to experience a He-flash but has a high enough mass for some $^3$He to survive. As a point of interest, we believe that K11 put a conservative lower limit on the mass range.  \\cite{eggleton08} modeled the destruction of $^3$He and found that as much as $\\sim 95$\\% of the $^3$He created during the MS evolution will be destroyed by the time the star reaches the tip of the RGB. For the 1.5 \\msun\\ model, the destruction fraction is only 75--83\\%. Nevertheless, the amount of $^3$He created in low mass stars is substantial. \\cite{iben67} estimates $\\sim10^{-3}$~\\msun, which corresponds to $A(^3{\\rm He})\\sim 8.6$. Assuming 95\\% of the original $^3$He was destroyed, only $10^{-4}$ of the remaining $^3$He must end up as $^7$Li to yield \\lli$=3.3$~dex (the abundance of our most Li-rich star). In the 1.5 \\msun\\ case, the fraction of $^3$He remaining is $\\sim20$~\\%, but the original $^3$He budget is smaller, so that the reservoir of $^3$He at the RGB tip only about 2.5 times larger in the 1.5~\\msun\\ case than in the 1~\\msun\\ case. K11 do not report \\vsini, so we cannot comment on any relationship of rapid rotation to \\lli\\ in their results.  However, K11 do provide \\cratio, and in contrast to  our stars, their warm Li-rich giants tend to have \\cratio$\\leq16$, indicating that some sort of extra mixing process has occurred in their Li-rich stars that does not appear to have occurred in the Li-rich stars studied here.    \\subsection{Planet Accretion} One of the main drawbacks to the Li-regeneration models is that they fail to explain the excess angular momentum of the rapid rotators. The accretion of a planet is one means by which a red giant star can acquire sufficient angular momentum to become a rapid rotator.  We can test whether our abundance results are consistent with the planet accretion paradigm by estimating the masses and chemical compositions  that  accreted planets would have had to account for  the mean abundance differences between the slow and rapid rotators. The expected  stellar abundances of Li after planet accretion---$A$(Li)$_{\\rm new}$---are  given by \\begin{equation} \\label{eq:liplanet} A({\\rm Li})_{\\rm new}= \\log(q_{\\rm e}10^{A({\\rm Li})_{\\rm p}}+10^{A({\\rm Li})\\star}) - \\log(1+q_{\\rm e}) , \\end{equation} where $q_{\\rm e}$ is the ratio of the planet mass ($M_{\\rm p}$) to the mass in the stellar convective envelope ($M_{\\rm env}$) and $A({\\rm Li})_{\\rm p}$ and $A({\\rm Li})_\\star$ are the initial abundances of Li in the planet and star, respectively.  (See the Appendix for the derivation of this equation.) For the present argument, we assume that the slow rotators are representative of the initial stellar abundances (i.e., $A({\\rm Li})_\\star =\\overline{A({\\rm Li})}_{\\rm slow}= -0.18$\\,dex) while the RV stable rapid rotators represent the post-planet-accretion  abundances (i.e., $A({\\rm Li})_{\\rm new} = 1.06$\\,dex).  If we assume that the meteoritic Li abundance  of our solar system  \\citep[\\lli~$\\sim$~3.3\\,dex,][]{lodders98} is a good representation of the Li abundances of hypothetically accreted planets (i.e, $A({\\rm Li})_{\\rm p}=3.3$),  then we find from Equation \\eqref{eq:liplanet} that  $q_{\\rm e}=5.4\\times10^{-3}$. This  $q_{\\rm e}$ corresponds to  $\\sim 6 M_{\\rm Jup}$  assuming $M_{\\rm env}=1$\\,\\msun, which is a reasonably-sized planet accreted into a reasonably-sized red giant convection zone.  Similarly, we can test our non-detection of a \\cratio\\ increase using our estimated  $q_{\\rm e}$. The carbon ratio expected after the accretion of a planet, $(^{12}{\\rm C}/^{13}{\\rm C})_{\\rm new}$, is given  by (see the Appendix) \\begin{equation} \\label{eq:cratio} (^{12}{\\rm C}/^{13}{\\rm C})_{\\rm new} = \\frac{10^{A({\\rm C})_{\\rm p}}\\frac{r_{\\rm p} q_{\\rm e}}{1+r_{\\rm p}}  + 10^{A({\\rm C})_\\star }\\frac{r_\\star }{1+r_\\star}}{10^{A({\\rm C})_{\\rm p}}\\frac{ q_{\\rm e}}{1+r_{\\rm p}} + 10^{A({\\rm C})_\\star }\\frac{1}{1+r_\\star}}, \\end{equation} where  $r_{\\rm p}=(^{12}{\\rm C}/^{13}{\\rm C})_{\\rm p}$ and  $r_\\star=(^{12}{\\rm C}/^{13}{\\rm C})_\\star$. For  a solar metallicity red giant star, we take $A({\\rm C})_\\star=8.26$, which is the solar-metallicity \\citep[$A$(C)$_\\sun=8.39$\\,dex; ][]{grevesse07} adjusted for some post-MS  processing of C and N such that [C/Fe]~$=-0.13$ \\citep{marcs08}. As with the lithium example, we use the average slow rotator \\cratio\\ to represent pre-planet accretion so that $r_\\star = 17.0$. Using Jupiter as an analog, we expect the assimilated planet to have had  $3\\times$  the  solar carbon abundance \\citep{wong04}. Thus, $A({\\rm C})_{\\rm p}=A({\\rm C})_\\sun+\\log(3)=8.87$\\,dex.  We adopt $r_{\\rm p} = 89$, which is a standard value in the solar system \\citep{lodders98}.  Under these assumptions and adopting  the $q_{e}$ derived from the Li enhancement,  we  calculate that  a sample of planet accreting stars (such as the rapid rotators) is expected to have $\\overline{^{12}{\\rm C}/^{13}{\\rm C}}=17.3$. Saturn is even more carbon-enriched \\citep{mousis09} with up to 10$\\times$ the solar value.  Using Saturn's composition as the analog for accreted planets, the expected carbon abundance increases to $\\overline{^{12}{\\rm C}/^{13}{\\rm C}}=18.0$. The average \\cratio\\ of the RV stable rapid rotators (17.3)  is similar to the value expected from accreting a Jupiter analog, supporting the plausibility of planet accretion by these stars. However, given the relative sizes of the signal (the \\cratio\\ increase) compared to our uncertainties, we caution that we cannot conclusively assert that we are measuring a true \\cratio\\ difference.   We can further test the plausibility of the planet accretion hypothesis by assessing whether both the enhanced angular momentum and  Li enrichment  are consistent. \\cite{carlberg09} introduced an equation (their Equation (1)) to relate an observed stellar \\vsini\\ to the properties of the star and the planet it accreted, i.e., $v\\sin i = 8[ M_{\\rm p}\\sin i \\sqrt{GM_\\star a_{\\rm p}(1-e^{2})} ]/M_{\\rm env}R_\\star$, where $M_{\\rm p}$ is the planet mass, $a_{\\rm p}$ is the planet's initial orbital separation, and $M_{\\rm env}$ is the mass in the stellar convection envelope. The terms in  brackets describe the initial angular momentum of a planet orbiting the star. To estimate the angular momentum gained by the star, both the approximate stellar envelope mass ($M_{\\rm env}$) and radius must be known  ($R_\\star$).   Because the latter can change by two orders of magnitude during RGB evolution, it is helpful to perform the analysis on individual stars as opposed to using the average stellar properties we have used thus far. To this end, we select a small group of Li-rich rapid rotators that represent the best candidates in our sample of stars that have accreted planets. In addition to the enhanced rotation and enriched Li, these are stars that have not yet evolved to the luminosity bump and are not suspected to be in close binary systems. Three stars meet all of these requirements and are listed in Table \\ref{tab:accretors}. The first two stars in that table are  the two most Li-rich stars in the ``Significant Dilution'' group, as seen in Figure \\ref{ali_by_group}. The third star, Tyc3340-01195-1, is the most Li-rich star  in the ``Dilution in Progress'' group and has a long-period binary companion. \\begin{deluxetable*}{lrrrrrr} \\tablewidth{\\textwidth} \\tablecaption{Best Case Candidates for Planet Accretion \\label{tab:accretors}} \\tablehead{ \\colhead{Star}&\\colhead{$T_{\\rm eff}'$}&\\colhead{\\vsini}&\\colhead{\\lli$_{\\rm LTE}$}&\\colhead{\\lli$_{\\rm NLTE}$}&\\colhead{\\cratio} &\\colhead{[Fe/H]}  \\\\ \\colhead{ }&\\colhead{(K)}&\\colhead{(\\kms)}&\\colhead{(dex)}&\\colhead{(dex)}&\\colhead{}&\\colhead{ (dex)} } \\startdata G0928+73.2600   & 4770   & 8.4  &    +3.62  &  +3.30 &     $28\\pm 8$ &$-0.26$ \\\\ Tyc0647-00254-1 & 4825   & 10.4  &   +1.92  &   +2.06 &    $20 \\pm 3$  &$-0.01$  \\\\ Tyc3340-01195-1 & 5040   & 8.4  &    +1.21  &   +1.32 &    $25 \\pm 5$ &$-0.18$  \\\\ \\enddata \\end{deluxetable*}   Matching the stars' effective temperatures, surface gravities, and metallicities to  grids of either stellar evolution tracks or isochrones can be used to estimate the stellar masses and radii. We opt to use isochrones in this analysis because the isochrones of \\cite{marigo08} can be interpolated onto a finer grid in $Z$ than what is available for the evolutionary tracks. We downloaded\\footnote{http://stev.oapd.inaf.it/cmd accessed on 2011 May 17.} isochrones with $Z=0.0002$, 0.001, 0.004, 0.008, 0.012, 0.016, 0.020, 0.024,  and 0.03 (Fe/H~= $-2.05$,  $-1.33$, $-0.72$, $ -0.41$, $-0.23$, $-0.10$, $0.00$, $+0.08$, and $+0.18$), and $\\log t$ ranging from 8.5 to 10.1 in increments of 0.1\\,dex (where $t$ is the age of the stellar population in years). The stellar masses and radii of the stars are estimated in the following manner.   First, we identify the isochrone metallicity that most closely matches the observed stellar metallicity.  For each age at constant $Z$, we find where the isochrone intersects the stellar \\teff. If there are multiple intersections, which is common for red giant evolutionary stages, we use the spectroscopically derived $\\log g$ to select the best intersection. The isochrones are computed on finite grids of \\teff\\  points, so we identify the isochrone points that straddle the stellar \\teff\\  and interpolate the masses and $\\log g$ associated with those two points to the \\teff\\ we are interested in. In other words,  for each star we reduce the isochrones to a grid of age, mass, and $\\log g$ for the stellar temperature. We then find adjacent grid points of isochrone-derived $\\log g$ that straddle the stellar $\\log g$,  and we interpolate the ages and masses associated with those points to estimate the age and mass of the stellar $\\log g$.   This last step generally finds between one and  three unique age/mass solutions, and the final mass and age estimates of our program stars  average over these unique solutions. Once we estimate the stellar mass, it is trivial to calculate an estimate in radius.  Using the measured surface gravity, the radius is given by $R_\\star=\\sqrt{GM_\\star/g}$, where $G$ is the gravitational constant. Although we acknowledge the large uncertainty inherent in this analysis, especially with the overlap of the RGB and horizontal branch at the stars' \\teff\\ and $\\log g$, we find that our isochrone fitting works rather well.  The sample of  giant SWPs has independent measurements  of mass  with which we can compare our results. In Figure \\ref{comp:massage}, we plot the comparison of our mass measurements to literature mass measurements and find that our masses agree with literature values within the quoted uncertainties, with the exception of Pollux, which has $M_{\\rm Lit.}=1.9$\\,\\msun\\ and $M_{\\rm isoch}=2.2\\pm0.2$\\,\\msun. \\begin{figure}[tb] \\centering \\includegraphics[width=1.03\\columnwidth]{f15.eps}   % \\caption{Comparison of the isochrone-derived masses of the giant SWPs ($y$-axis) with the mass measurements reported in the Exoplanet Encyclopedia ($x$-axis). A unity-slope line is plotted for reference. \\label{comp:massage}  } \\end{figure}   The isochrones do not provide information on the internal structure of the stars. Therefore, we use the \\cite{giard00} stellar evolution models that most closely match the mass and metallicity of our stars to estimate the fraction of the mass in the convective envelope.  The least evolved star in Table \\ref{tab:accretors} is Tyc3340-01195-1,  which has $\\sim32$\\% of the total mass  in the stellar envelope.  The other two stars have envelope mass fractions of $\\sim 80$\\%. \\begin{figure}[tb] \\centering \\includegraphics[width=1.03\\columnwidth]{f16.eps}   % \\caption{Minimum ``projected mass'' ($M_{\\rm p} \\sin i$) needed to account for the observed \\vsini\\ of the best case planet accretion candidates as a function of the expected maximum initial semimajor axis  ($a_0$) of the accreted planet. The true mass of any accreted planets will be larger both for smaller inclination angles (more pole-on angles) and for $a_0<{\\rm max.}(a_0)$. For reference, the planetary masses and orbital separations of known extrasolar planets orbiting MS stars  are shown with crosses. \\label{fig:pacc}  } \\end{figure}   Obviously, the unknown inclination angle ($i$) limits us to estimating the minimum planet mass ($M_{\\rm p}\\sin i$). There is also a degeneracy between the planetary mass and orbital separation.  A massive planet at a small separation can create the same degree of stellar rotation as a less massive planet at an initially larger separation. Fortunately, the tidal modeling of \\cite{carlberg09} provides a means of breaking this degeneracy in the limiting case of the maximum initial orbital separations of planets that could have been accreted ($a_{\\rm max}$)---planets that are too distant from their stars will not be accreted.  In the  exoplanetary systems  modeled in \\citet[][99 systems with a total of 115 planets]{carlberg09}, we find that the ratio of the planets' initial semimajor axis ($a_{\\rm 0}$) to the stellar radius at the time of planet accretion, $R_\\star(t_{\\rm acc})$,  is  typically 4.2. Therefore,  we can  estimate $a_{\\rm max}\\sim 4.2\\, R_\\star$.   Substituting this maximum orbital separation into the angular momentum equation yields the minimum planet mass capable of producing the observed rotational velocities in the rapid rotators.   This derived mass is both a minimum $M_{\\rm p}$ because  of the unknown inclination {\\it and} a minimum $M_{\\rm p}\\sin i$ because more massive planets with smaller initial $a_{\\rm p}$ are also capable of producing the observed rotation. The results of this calculation for our selected Li-rich rapid rotators   are presented in Figure \\ref{fig:pacc}, which shows the minimum $M_{\\rm p}\\sin i$ required to spin up the  star as a function of $a_{\\rm max}\\sim 4.2 R_\\star$.    The least evolved star could have accreted planets that originally orbited within 0.1~AU and requires a Jupiter-mass planet to account for the angular momentum.   The other two red giants could have accreted planets within 0.2~AU, and the planets must have been at least 4~$M_{\\rm Jup}$ to explain their enhanced rotation. This minimum planet mass is  slightly less than the average of 6~$M_{\\rm Jup}$ planet needed to account for the Li-enrichment in Section \\ref{sec:meanlicarbon}.   For comparison, we also plot in Figure \\ref{fig:pacc} the planetary masses and orbital separations of known exoplanets.\\footnote{Data come from the Exoplanet Orbit Database \\cite{wright11}, accessed on  2011 December 18.} Our angular momentum analysis implies that our rapid rotators can be explained by accreting planets to the left and above the  rapid rotators' positions on the plot, and we find such planet do exist among the known exoplanetary systems.    However, it is also worth noting that there is a dearth of Jupiter mass planets (and larger) between $\\sim 0.08$ and $\\sim 0.7$~AU, the region where stars like those listed in Table \\ref{tab:accretors}  would be actively  clearing out their planets by tidal decay.    \\subsection{A Few Noteworthy Stars} {\\it G0928+73.2600.} This star was analyzed separately  by \\cite{carlberg10b}, with attention brought to its apparent pre-bump evolutionary stage but high enough \\lli\\ to require accreting a Li-enriched object. We highlight this star again to contrast it to another star in the study that has almost identical stellar parameters, but very different \\lli\\ and \\cratio. Using the stellar parameters from the 2007 observations, G0928+43.2600 has \\teff$=4900$~K, \\logg$=2.70$, and [Fe/H]$=-0.26$, and the slow rotator G1200+67.3882 has \\teff$=4900$~K, \\logg$=2.70$, and [Fe/H]$=-0.24$.   Aside from rotation speeds (8.4~\\kms\\ and 1.5~\\kms), these stars differ in their abundances. The rapid rotator has \\lli$=3.30$ and \\cratio$=28$, while the slow rotator has \\lli$=+0.39$ and \\cratio$=8.6$.  Such a low \\cratio\\ is suggestive of more than the usual mixing, which implies that this star could have regenerated Li internally and is now in the process of destroying that regenerated Li and reducing  \\cratio\\ below the normal values. If this scenario is indeed true,  it implies that G0928+73.2600 is at the slightly earlier stage where newly regenerated Li has not yet been destroyed and so the \\cratio\\ is still at the standard value.  The problem with using Li regeneration to explain the abundance differences of these otherwise similar stars is that the stars are both too hot to have evolved through the bump phase  ($T_{\\rm bump}\\sim $~4600--4700~K for [Fe/H]=$-0.25$). This difficulty could be resolved if our choice of stellar evolution tracks predicts too cool temperatures for the luminosity bump or if our derived spectroscopic temperatures are systematically too hot. To test the first scenario, we compare stellar evolution tracks from two independent sources. \\cite{charbonnel00} plot solar metallicity stellar evolution tracks in their Figure 1. A visual inspection of that plot shows the luminosity bump at temperatures between 4450 and 4500~K, which is $\\sim 100$~K cooler than the luminosity bump of the solar-metallicity models we plot in Figure \\ref{model_dilution}. As a second test, we computed the evolution of a 2~\\msun, $Z=0.011$ star using the MESA code \\citep{paxton11}\\footnote{MESA version 3661, using the ``inlist'' file provided in the  ``1M\\_pre\\_ms\\_to\\_wd''  test suite as a template. We changed only the mass and metallicity of that file.}.  The luminosity bump of that model extends to a temperature of 4700~K, comparable to the models shown in Figure \\ref{model_dilution}. Next, we compile two different photometric temperatures to test for a systematic offset in our spectroscopic derivation. The first set of temperatures comes from \\cite{carlberg11} and were derived  using the stars' Washington $M-T_2$ colors \\citep[which are converted to Cousin $V-I$, ][]{majewski00} and \\cite{houd00} color-temperature relations.  This calculation yields $T_{\\rm phot1}= 4773$~K for the rapid rotator and $T_{\\rm phot1}= 4875$~K for the slow rotator. Both stars also have {\\it Tycho} designations.   G0928+73.2600 is Tyc4382-00780-1, and G1200+67.3882 is Tyc4160-00999-1.  We converted their observed {\\it Tycho} $B-V$ magnitudes \\citep{hoeg00} to Johnson $B-V=0.85(B-V)_{\\rm Tycho}$, dereddened the colors using \\cite{schlegel98} maps, and again used \\cite{houd00} color-temperature relations to find  $T_{\\rm phot2}= 4755$~K for the rapid rotator and $T_{\\rm phot2}= 5002$~K for the slow rotator. Thus, the rapid rotator may be slightly cooler and more evolved than our spectroscopic analysis suggested, while the slow rotator may be slightly warmer (and less evolved).  However, the photometric temperatures are still warmer than the luminosity bump temperatures. Therefore, these two stars may represent  examples of the He-flash lithium regeneration hypothesized by K11.  {\\it Tyc3340-01195-1.} This star is a long period binary star, and C11 argued that the  relatively large separation of the stellar components in the former system  ($a_\\star\\sin i\\sim 425 R_\\sun\\sim 2$~AU)    made tidal synchronization an unlikely  explanation for the primary star's enhanced rotational velocity  (\\vsini=8.4\\,\\kms).  The presence of a stellar companion raises  questions about the stability of planets in the system. However,  planets can have stable orbits around the primary star  interior to the  stellar companion  if the orbits are small enough. In a study of such ``S-type'' planetary orbits, \\cite{rabl88} found a quadratic relationship relating the largest Lowest Critical Orbit (LCO---orbits larger than this may be unstable, while smaller orbits are stable) to the binary system's separation ($a_\\star$) and eccentricity ($e$) to be ${\\rm LCO}=0.262a_\\star-0.254a_\\star e-0.060a_\\star e^2$.  For the Tyc3340-01195-1 system, a planet orbiting with $a_{\\rm p}<0.4$~AU would be stable around the primary. In Figure \\ref{fig:pacc} we find that the maximum initial separation of planets that could have been accreted by Tyc3340-01195-1 is only $\\sim0.1$~AU; therefore, the orbits of any accreted planets would have been dynamically stable in the binary star system during the primary star's MS lifetime.     We have studied the global abundance patterns of \\lli\\ and \\cratio\\ in a homogeneously selected sample of slow and rapid rotators to see whether  the rotation is correlated with the replenishment of elements destroyed during the stellar evolution, as expected from planet accretion. Our final sample contains 71 slow rotator stars (58 chosen purely for their slow rotation, 10 selected because they are known to host planets, and the three M08 stars) and 15 rapid rotators. From an analysis of line-of-sight considerations, we expect that 4.5\\% of our slow rotators have true rotational velocities qualifying as rapid rotation.   Overall, the rapid rotators show an enhancement of \\lli\\ over the slow rotators by 0.99\\,dex. When selecting the subset of rapid rotators that are RV stable, the average enhancement of \\lli\\   increases to 1.24\\,dex over the slow rotators.  These Li enhancements are  consistent with the accretion of a $\\sim 6$ Jupiter masses of planetary material with a Li abundance similar to meteoritic abundances (thought to be relic of the solar nebular abundance).  Consistent with this explanation and our relatively large \\cratio\\ uncertainties, we find no statistically significant difference between the \\cratio\\  measured in the rapid and slow rotator samples. A more massive or more carbon-enriched object would have to be accreted to be measurable at the level of our \\cratio\\ uncertainties. However, we do not measure smaller \\cratio\\ in our Li-rich stars or rapid rotators, a signature that  would  need to be seen to prove  the replenishment of Li through nuclear processing. We also compared our stellar sample to evolutionary tracks to ascertain how the relative Li and carbon abundances varied within subsets of similar evolutionary stage.  This comparison was necessary because the rapid rotators  and slow rotators are not distributed evenly across the RGB.   We found that in all groups that contained more than one rapid rotator, the most Li-rich star in the group was a rapid rotator.  In other words, the result that  rapid rotators  are more Li-rich than the slow rotators persists even within groups of similarly-evolved stars.  These main conclusions were drawn by comparing the global properties of our two main samples, where it is safe to assume that the peculiar properties of individual stars are  likely to average out. However,  to determine whether both the Li enrichment and enhanced angular momentum were consistent with the planet accretion hypothesis, we needed estimates of both the stellar mass and radius. The latter can vary by two orders of magnitude during RGB evolution.  Instead of averaging these properties over our stellar sample, we selected three Li-rich rapid rotators  that are the best examples in our sample of stars that may have engulfed planets. Stellar isochrones were used to estimate stellar masses and radii.  Combining these stellar properties with the stars' measured \\vsini\\ and an estimate of the maximum orbital separations of accreted planets, we find that minimum planet masses of $\\sim4.5$~$M_{\\rm Jup}$ can account for the rotation of the two most Li-rich, rapidly rotating stars. This planetary mass estimate is comparable to the $\\sim 6$~$M_{\\rm Jup}$ needed to explain  the  global Li enrichment seen in the rapid rotators. The rotation of the third (and least Li-rich) of the selected stars can be explained with a minimum planet mass of only $1$~$M_{\\rm Jup}$.   However, we also found in the stellar evolution analysis that our stellar sample did not reproduce the detailed \\lli\\ and \\cratio\\ abundance patterns we expected. First, there is significant scatter in the \\lli\\ abundances in every evolution group, and we found stars with low \\cratio\\ at stages earlier than the completion of FDU. This scatter may be due to the wide range of stellar metallicities and masses represented in our sample combined with both a temperature-dependent \\lli\\ sensitivity and a difficulty in measuring precise values for \\cratio~$\\gtrsim 20$. Second, we did not find any classical Li-rich stars (\\lli\\ $\\gtrsim 1.5$~dex) at the luminosity bump. The most Li-rich slow rotator had \\lli$=1.12$ (and may be associated with the luminosity bump, the fourth of our six evolutionary classes). Furthermore, all stars that are more Li rich than this  are in the second of our defined classes---likely at pre-bump evolutionary stages.   This finding is reminiscent of the \\cite{kumar11} study that also found Li-rich giants at pre-bump evolutionary stages.   Based only on this concentration in \\teff, they suggested that the stars in their sample are red clump stars and that Li regeneration may occur during the He flash.  However, most of the warm (\\teff$\\leq 4600$\\,K), pre-bump Li-rich stars in the K11 study had \\cratio$\\leq 16$, whereas our Li-rich pre-bump stars generally have \\cratio$\\sim 20$. On the other hand, we noted that there is a slow rotator in our sample with nearly identical \\teff, \\logg, and [Fe/H] as the most Li-rich star in our sample. That slow rotator has a low \\cratio\\ suggestive of the enhanced extra mixing that should only occur at post-luminosity bump stages. Together, these two stars appear to be in adjacent stages of the Li-regeneration phenomena that is thought to occur at they luminosity bump except that they are both too hot to be luminosity bump stars.  In conclusion, the \\lli\\ and \\cratio\\ of our sample showed far greater complexity than we anticipated.  The tendency for the rapid rotators  to show Li-enrichment implies that either planet accretion or some sort of rotational mixing has taken place in these stars.   If planet accretion is not responsible for the Li enriched stars, then the fact that these stars are hotter than the luminosity bump presents a problem for the Li-regeneration models that generally require the removal of the mean molecular weight barrier. Many of the slow rotators in our sample showed lower than expected \\lli, suggesting that a variety in either the initial Li abundances or amount of Li destruction on the MS may exist."
    },
    "1208/1208.4366_arXiv.txt": {
        "abstract": "We combine star-formation histories derived from observations of high redshift galaxies with measurements of the $z\\sim 0$ relation between gas-phase metallicity, stellar mass, and star formation rate to make an explicit and completely empirical connection between near-field and distant galaxy observations. Our approach relies on two basic assumptions: 1) galaxies' average paths through time in stellar mass vs.\\ star formation rate space are represented by a family of smooth functions that are determined by the galaxies' final stellar mass, and 2) galaxies grow and become enriched with heavy elements such that they always evolve along the mass--metallicity--star formation rate relation.  By integrating over these paths, we can track the chemical evolution of stars in galaxies in a model independent way, without the need for explicit assumptions about gas inflow, outflow, or star formation efficiency. Using this approach, we present predictions of stellar metallicity (i.e., O/H) distribution functions for present day star-forming galaxies of different stellar masses and the evolution of the $\\alpha$-element stellar metallicity-mass relation since $z\\sim 1$. The metallicity distribution functions are fairly well described as Gaussians, truncated at high metallicity, with power-law tails to low metallicity. We find that the stellar metallicity distribution for Milky Way mass galaxies is in reasonable agreement with observations for our Galaxy, and that the predicted stellar mass vs. mean stellar metallicity relation at $z=0$ agrees quite well with results derived from galaxy surveys. This validates the assumptions that are implicit in our simple approach. Upcoming observations will further test these assumptions and their range of validity, by measuring the mean stellar mass-metallicity relation up to $z\\sim 1$, and by measuring the stellar metallicity distributions over a range of galaxy masses. ",
        "introduction": "\\label{sec:intro} As stars are born with the imprint of the elemental abundances in their gaseous birth environments, measurements of stellar metallicities within galaxies can be used to gain insight into their star formation and chemical histories \\citep{tinsley75,tinsley78,blandhawthorn10}. In this paper, we carry out an explicit test of the self-consistency between empirically derived star formation histories, observational scaling relations between gas phase metallicity ($\\zg$), stellar mass ($\\mstar$), and star formation rate ($\\msfr$), and observed stellar metallicities at $z=0$.  It has long been known that there is a strong positive correlation between the luminosity or stellar mass of star forming galaxies and their galaxy-averaged gas-phase abundance \\citep{garnett87,zaritsky94}.  This relationship (the ``\\mzr'') has now been measured for large samples of nearby galaxies \\citep[e.g.][]{tremonti04,kobulnicky04,kewley08} and for smaller samples of distant galaxies out to $z\\sim 3.5$ \\citep{savaglio05,shapley05,erb06a,maiolino08,mannucci09}.  The normalization, slope, scatter, and evolution of the observed \\mzr\\ place strong constraints on theoretical models of galaxy formation, in particular on the physics of feedback from massive stars and supernovae. Recently, it has been suggested that the scatter in the observed \\mzr\\ can be reduced by considering the star formation rate as a third parameter \\citep{laralopez10,mannucci10}.  Physically, the star formation rate is probably acting as a proxy for the galaxy gas fractions, as galaxies with higher gas fractions have both higher star formation rates and more diluted metals \\citep{hughes12}.  Moreover, \\citeauthor{laralopez10} and \\citeauthor{mannucci10} find that galaxies up to $z\\sim2.5$ \\citep[e.g.,][]{savaglio05,erb06a} appear to always lie on the same $\\mstar$-$\\zg$-$\\msfr$ ``fundamental'' relation, though this may begin to break down at $z\\sim 3$ (\\citealp{maiolino08}, though see also \\citealp{yabe12}).  It is not yet clear whether or why a non-evolving mass--metallicity--star-formation rate relation should be generically expected, although some cosmological models do predict this (\\citealp{dave12}; M.\\ Arrigoni et al.\\ in preparation). Self-regulation due to stellar-driven outflows is usually invoked in this context \\citep{finlator08,spitoni10,peeples11,dave11,dayal12}.  There is a long history in the literature of attempts to model the chemical evolution of galaxies using a ``classical'' approach, the simplest version of which is the ``closed box'', in which galaxies start out with all of their mass in the form of pristine gas, and convert some fraction of that gas to stars, producing a certain ``yield'' of heavy elements along the way. It was quickly realized that the simplest closed box picture could not reproduce observations, in particular of the distribution of stellar metallicities in the Solar Neighborhood in our Galaxy (the ``G-dwarf'' problem), leading to modifications such as inflowing pristine or pre-enriched gas as well as outflows of gas and metals \\citep{lyndenbell75,pagel89,colavitti08}. While much has been learned from this approach, it has the obvious drawback that the results tend to be dependent on the rather arbitrary functions adopted for the inflow and outflow rates.  Another approach that has become fairly widely used is the coupling of detailed chemical evolution models with cosmological models of galaxy formation, realized in the form of either numerical hydrodynamic simulations \\citep{scannapieco06,kobayashi07,brooks07,mouhcine08,oppenheimer08,wiersma09} or semi-analytic models \\citep{delucia04,nagashima05,somerville08,arrigoni10}. These have the advantage that quantities such as the rate of inflow of gas into galaxies is motivated by the formation rate of structure in the Cold Dark Matter (CDM) paradigm that provides the backbone for these models, and outflow rates are also governed by physically motivated scalings. However, they suffer from the drawback that currently, all CDM-based galaxy formation simulations, whether numerical or semi-analytic, apparently fail to reproduce the observed scaling of star formation history functional form with galaxy mass (sometimes referred to as ``downsizing''; see \\citealp{fontanot09} for a detailed discussion).  Here we propose a different approach to modeling the buildup of metals in galaxies that sidesteps some of the problems with both the classical and cosmological methods. We make use of a family of parameterized empirical star formation histories, derived from observations of the stellar mass vs.\\ star formation rate relation at different redshifts by \\citet{leitner12}. These empirically-derived star formation histories are in good agreement with those derived from spectral energy distribution modeling and the evolution of the stellar-halo mass relation \\citep{zheng07,conroy09a,behroozi12,moster12}.  We assume that the {\\em average} star formation history for galaxies with a given stellar mass today can be represented by this smooth functional form. This assumption is supported by mounting observational evidence that galaxies with masses of roughly the Milky Way or lower ($\\sim$ few $\\times 10^{10}\\, \\msun$) build up most of their mass through smooth accretion of gas, with mergers and starbursts playing a relatively minor role \\citep{noeske07a,robaina09,oliver10,karim11,leitner12}. This is also supported by results from cosmological simulations \\citep{brooks09,hirschmann12,tissera12}. Then, once stellar mass loss and recycling has been accurately accounted for \\citep{leitner11}, the growth of galaxies can be traced through cosmic time by stepping along these $\\msfr(\\mstar)$ relations as a function of redshift.  We make a second assumption, that galaxies always evolve {\\em along} the $\\mstar$-$\\zg$-$\\msfr$ relation. That means that as galaxies evolve along their trajectories in $\\mstar$-$\\msfr$ space (specified by their star formation history), we assume that the stars born at that moment form out of gas with a metallicity $\\zg$ given by the $\\mstar$-$\\zg$-$\\msfr$ relation. In this way, we can build up the {\\em distribution function} of stellar metallicities for galaxies with different masses at $z=0$. We can also test whether the two assumptions of empirical star formation histories plus a non-evolving $\\mstar$-$\\zg$-$\\msfr$ relation lead to an average stellar metallicity relation at $z=0$ that agrees with observations. Both of these quantities will be better constrained observationally over a wider range of galaxy masses, and to higher redshift, in the near future.  This paper is organized as follows.  In \\S\\,\\ref{sec:relations}, we describe our adopted and predicted scaling relations connecting galaxy stellar masses, star formation rates, and gas-phase metallicities across cosmic time.  In \\S\\,\\ref{sec:stars} we present the stellar metallicity distribution functions and evolution of the stellar metallicity-mass relation predicted by this model.  We conclude with a brief discussion of relevant implications in \\S\\,\\ref{sec:conc}.  Throughout we assume a \\citet{chabrier03b} stellar initial mass function (integrated over $0.1$--$100\\,\\msun$) and \\tlogoh$_{\\odot}=8.7$ \\citep{asplund09}, although we note that the 8.9 value found from helioseismology \\citep{delahaye06} would lead to $0.2$\\,dex lower $Z/\\zsun$ values than stated here.  Following \\citet{leitner12}, we adopt a flat $\\Lambda$CDM cosmology with $\\Omega_{\\rm M}=1-\\Omega_{\\Lambda}=0.258$ and $h=0.72$. ",
        "conclusions": "\\label{sec:conc} We have combined empirically derived star-formation histories from observations of high redshift galaxies \\citep{leitner12} with the empirically determined relation between stellar mass, star formation rate, and gas-phase metallicity \\citep{mannucci10} to derive the distribution of stellar metallicities at $z=0$ and the evolution of the galaxy-averaged stellar metallicity-mass relation (Figure\\,\\ref{fig:stellar} and Tables\\,\\ref{tbl:hist} and \\ref{tbl:mzevolve}).  We find that the hypothesis that the observed $z=0$ relation between stellar mass, star formation rate, and gas-phase metallicity holds up to at least $z\\sim2.5$ is consistent with stellar metallicities, within the limits of the current observations.  This comparison, however, is hampered by both a lack of a well-measured $\\mstar$-$\\msfr$-$\\zg$ relation at low stellar masses, low star formation rates, and high star formation rates, and a lack of well-characterized $\\alpha$-element distribution functions in both the Milky Way and other star-forming field galaxies with $10^8\\lesssim\\mstar\\lesssim 10^{11}\\msun$. From the modeling side, upcoming characterizations of the $\\mstar$-$\\msfr$-$\\zg$ relation to lower stellar masses and across a broader range of star formation rates (B.\\ Andrews \\& P.\\ Martini, in preparation) will help alleviate the need for extrapolation when calculating \\tlogoh.  Upcoming surveys will also provide a more complete census of the stellar metallicity distribution function for the Milky Way (e.g., the SDSS-III APOGEE survey) and Local Group dwarf galaxies (E.\\ Kirby et al.\\ in preparation).  We have assumed here that all of the scatter in \\tlogoh\\ at fixed stellar mass is related to variations in the star formation rate, but the resulting relation still has scatter that is correlated with, e.g., galaxy size \\citep{ellison08a,yabe12} or environment \\citep{kewley06a,ellison08b,pasquali12,hughes12}.  Moreover, the scatter in $\\zstar$ at fixed stellar mass is known to correlate with galaxy age \\citep{gallazzi05}.  In our framework, this can be seen as variations in galaxies' star formation histories via the $\\sim0.3$\\,dex\\ scatter in star formation rate at fixed stellar mass \\citep{noeske07a,karim11,leitner12}.  This variation can be partially attributed to changes in environment \\citep{pasquali10}, especially if environment introduces additional scatter to the $\\zg$-$\\mstar$-$\\msfr$\\ relation \\citep[e.g.,][]{pasquali12}.  We have also assumed that the buildup of {\\em all} of the stellar mass in a galaxy can be well described by Equation\\,(\\ref{eqn:sfms}).  This is clearly not the case, as, e.g., there are stars in the Milky Way that were formed before $z\\sim2.5$ and with $\\log \\zstar/\\zsun < -3$. However, these stars comprise a very small fraction of the total stellar mass or light, at least in the Milky Way, and could be addressed by more complete model of star formation histories, including minor mergers which bring in a populations of lower-metallicity stars.  One of the strengths of the method presented here is that it does not require us to explain {\\em why} the relationship between gas-phase metallicity, stellar mass, and star formation rate apparently does not evolve with redshift by invoking carefully evolving gas fractions, accretion rates, or outflow efficiencies \\citep[such as in, e.g.,][]{peeples11,dave12,dayal12}.  On the other hand, integrating over the evolution of the \\mzr\\ does neatly predict what fraction of Oxygen ever produced by a $z=0$ galaxy should be still locked up in stars, and how steeply this fraction should increase with stellar mass \\citep[c.f.,][]{gallazzi08,kirby11c,zahid12b}.  Moreover, by combining this kind of calculation with measurements of the \\mzr\\ and $z=0$\\ gas fractions \\citep[e.g.,][]{peeples11}, one can place constraints on the total amount of Oxygen galaxies have expelled over their lifetime \\citep{bouche07,kirby11c,zahid12b}, though for star-forming galaxies, uncertainties in the gas-phase abundances and relevant gas masses make this calculation difficult. More globally, it is especially intriguing that the implied masses of both metals and all baryons expelled by galaxies over their lifetime appears potentially consistent with the oxygen mass observed in the circumgalactic medium of $z\\sim 0.25$ galaxies (\\citealp{tumlinson11}; J.\\ Werk et al.\\ in preparation)."
    },
    "1208/1208.4685_arXiv.txt": {
        "abstract": "We show that the existence of { prograde} equatorial satellites is { consistent} with { a} collisional tilting scenario for Uranus. In fact, if the planet was surrounded by a proto-satellite disk at the time of the tilting and a massive ring of material was temporarily placed inside the Roche radius of the planet by the collision, the proto-satellite disk would have started to precess incoherently around the equator of the planet{, up to a distance greater than that of Oberon}.  Collisional damping would then have collapsed it into a thin equatorial disk, from which the satellites eventually formed. The fact that the { orbits of the} satellites are prograde requires Uranus to have had a non-negligible initial obliquity (comparable to that of Neptune) { before} it was finally tilted to 98 degrees. ",
        "introduction": "The origin of the large obliquity of Uranus remains elusive. Two scenarios have been proposed: an impulsive tilt due to a collision with a massive body (Safronov, 1966) or a slow tilt due to a resonance between the precession rates of the spin axis and of the orbit (Bou\\'e and Laskar, 2010).  A critical constraint, inherent to both of these scenarios, is that the regular satellites of Uranus have essentially equatorial orbits and are prograde relative to the rotation of the planet. Notice that the rotation of the planet is, strictly speaking, retrograde, as Uranus obliquity is about 98 degrees.  In principle, if the satellites were originally coplanar to the equator of Uranus and the planet was tilted slowly (as in Bou\\'e and Laskar, 2010), the satellites would have preserved equatorial orbits by adiabatic invariance. Indeed, in the system's current configuration, the Laplace plane ({the} reference plane about which satellite orbits precess) is very close to Uranus' equatorial plane, for all bodies up to Oberon's distance, due to the oblateness of the planet.   In order to tilt Uranus slowly, a resonance between the precession rates of the spin axis and of the orbital plane is required. This means that the former had to be much faster than it is today. In Bou\\'e and Laskar (2010) this is achieved by assuming that Uranus originally had a massive satellite with an orbital radius of about 0.01 AU (Satellite X, hereafter). This assumption, however, is problematic because the Laplace plane for Satellite X is close to the orbital plane of Uranus. Thus, Satellite X would not follow the equator during the tilting of the planet and, by virtue of its large mass, would retain the other satellites (particularly Titania and Oberon, as we verified by numerical integration of a slowly tilting system) { near} its own orbital plane.  When Satellite X is removed by chaotic dynamics, the tilting process is over. Yet, the orbits of the regular satellites of Uranus would remain off equator, as in the impulsive tilting scenario.  In absence of a slow-tilting scenario that does not invoke the existence of Satellite X, we are left with the impulsive tilting scenario as the only viable option. Accordingly, in this Note, we investigate the conditions under which the equatorial, prograde orbits of the regular satellites of Uranus can be reproduced in the context of the collisional tilting scenario.  {For clarity, we proceed in steps.  We first investigate in Sections~2\\&3 the dynamics of a proto-satellite disk around an oblated planet with a tilted spin axis. This highlights the competing effects of the planet's $J_2$, the solar perturbation and the self-gravity of the disk. We focus on the conditions that lead the disk to precess incoherently around the planet's equatorial plane and, eventually, to collapse into an equatorial disk. We don't worry at this stage on whether this equatorial disk would be prograde or retrograde. Thus, although we fix the obliquity at the current value for Uranus, the dynamics that we study are basically independent on the obliquity value. Then, in Section~4 we consider the fact that Uranus is a retrograde planet (obliquity $\\epsilon=98^\\circ$). Thus, a disk originally on the orbital plane of Uranus (as expected if Uranus was tilted from $\\epsilon=0^\\circ$ to $98^\\circ$ in one shot) would necessarily become an equatorial, {\\it retrograde} disk. We then investigate which tilting histories of the planet would have a non-negligible probabilities to produce a {\\it prograde} equatorial disk. In Section 5, we finally discuss the implications of these tilting histories on our understanding of giant planets growth.} ",
        "conclusions": "We conclude that the collisional tilting scenario for Uranus is consistent with the prograde, equatorial character of the orbits of its regular satellites, { as well as the size of the system}. The fact that the satellites are prograde, implies that Uranus was not tilted from 0 to 98 degrees in one shot. Instead, it had to have { had} a non-negligible obliquity, { prior to} the { final} giant impact. Thus, { Uranus} should have experienced at least two giant collisions.  This result, together with the obliquity of Neptune, which also has no other explanation than a collisional tilt, suggests that giant impacts, affecting the obliquities, { were} rather common during the growth of the ice-giants of the solar system. {Thus, these planets presumably did not grow by the sole accretion of small planetesimals as often envisioned.  Instead, towards the end of their accretion history, they should have experienced a phase similar to that characterising the process of terrestrial planet formation, i.e. dominated by giant impacts with other large planetary embryos. Past and future models of growth of giant planet cores should be confronted with this constraint.}"
    },
    "1208/1208.2295_arXiv.txt": {
        "abstract": "{Our knowledge on the central components of disk galaxies has grown substantially in the past few decades, particularly so in the last. This frantic activity and the complexity of the subject promote confusion in the community. In these notes, I discuss the concept of galactic bulge and its different flavors. I also address fundamental scaling relations and the bulge-elliptical galaxy connection, their central black holes and formation models. In particular, I aim at conveying three important notions: {\\bf (i)}: box/peanuts are just the inner parts of bars; {\\bf (ii)}: the physical reality of two different families of bulges is evident; and {\\bf (iii)}: at the high mass end, at least, classical bulges are {\\em not} just scaled down ellipticals surrounded by disks. ",
        "introduction": "\\begin{figure*} \\begin{center} \\resizebox{0.15\\hsize}{!}{\\includegraphics[clip=true]{date1.eps}} \\hspace{0.35\\hsize} \\resizebox{0.15\\hsize}{!}{\\includegraphics[clip=true]{date2.eps}}\\\\ \\end{center} \\vskip-0.25cm \\resizebox{0.5\\hsize}{!}{\\includegraphics[clip=true]{auth70-00.eps}} \\resizebox{0.5\\hsize}{!}{\\includegraphics[clip=true]{auth01-.eps}}\\\\ \\begin{center} \\vskip-1cm \\resizebox{0.35\\hsize}{!}{\\includegraphics[clip=true]{abs70-00.eps}} \\hspace{0.15\\hsize} \\resizebox{0.35\\hsize}{!}{\\includegraphics[clip=true]{abs01-.eps}} \\end{center} \\caption{ {\\footnotesize Top left: authors with first-author papers in the period 1970--2000 in ApJ, AJ, MNRAS and A\\&A with the words `bulge(s)' and `galaxy(ies)' in the abstract. The word-cloud is limited to authors with three or more publications, and the font size is proportional to the number of papers. Bottom left: most common relevant words in the abstracts of all such publications in the period. Top right and bottom right are the corresponding word clouds for the period 2001--2012 (mid-May). In the 1970-2000 period, this search returns 1562 published papers, and 143 authors with more than three first-author papers. For the period 2001--2012, these figures change to 1999 papers and 178 authors.}} \\label{fig:lit} \\end{figure*}  These notes correspond to a couple of Lectures given at the School of Astrophysics ``F. Lucchin'' for PhD students and young researchers, held in Erice, Italy, in September 2011. One of the two subjects of the School was Galaxy Bulges, and the presentation slides are available online\\footnote{See \\href{http://www.sc.eso.org/~dgadotti/astro.html}{http://www.sc.eso.org/$\\sim$dgadotti/astro.html}.}. The content in the slides is significantly more extended than what the limited space here allows, and I stay considerably on the deceptively simple, difficult subject of bulge definitions. Current literature abounds with confusion, and I thus dedicate space to try and shed some light on this topic, not only for those beginning their way, but hopefully also for a broader audience in need.  I would like to right away acknowledge reference publications which have influenced my view substantially. These are \\citet{BinTre87}, \\citet{WysGilFra97}, \\citet{BinMer98}, \\citet{KorKen04} and \\citet{Ath05b}. Also important are the relatively recent Conference Proceedings of the IAU Symp. 245, and the recent review by \\citet{Gra11}. Although I did my best to cope with the enormous body of literature covering the subject, the reference list is but a tiny fraction of it. In order to minimize this inherent bias in these Notes, Fig. \\ref{fig:lit} displays word-clouds with the first authors of papers on galaxy bulges published in two different periods: 1970--2000 and 2001--2012 (mid-May). Font sizes are porportional to number of papers, rather than citations, as the latter are also biased to some extent. I hope that this will alert the reader to authors and studies other than those I quote here. Figure \\ref{fig:lit} also shows word-clouds made with common words in the abstracts of these publications. It is interesting to see that these words have not changed much in the two periods, with few notable exceptions, including the word `black-hole'. ",
        "conclusions": ""
    },
    "1208/1208.0635_arXiv.txt": {
        "abstract": "Shadows of multi-black holes have structures distinct from the mere superposition of the shadow of a single black hole:  the eyebrow-like structures outside the main shadows and the deformation of the shadows. We present analytic estimates of these structures using the static multi-black hole solution (Majumdar-Papapetrou solution). We show that the width of the eyebrow is related with the distance between the black holes and that the shadows are deformed into ellipses due to the presence  of the second black holes. These results are helpful to understand qualitatively the features of the shadows of colliding black holes. We also present the shadows of colliding/coalescing black holes in the Kastor-Traschen solution. ",
        "introduction": "One of the hottest topics in galaxy formation is the coevolution of super massive black holes (SMBHs) with spheroid components (bulges) of galaxies. It becomes more and more clear that most of galaxies and AGNs have at least one SMBH and there is a strong correlation between the SMBH mass and the bulge mass of host galaxies~\\cite{kormendy,magorrian,merritt}. Although a detailed mechanism of the coevolution is not yet understood, it is almost certain galaxy mergers play an essential role, since it is known that bulges or spheroid components are formed by merger of galaxies in the hierarchical clustering scenario of structure formation. Hence it may be also natural to consider formation of SMBHs is due to mergers of smaller black holes.  An observation clue of existence of binary black holes is recently obtained from detailed study of Kepler motion of a radio emission component in the radio galaxy 3C 66B by using a technique of phase-referencing very-long-baseline interferometry (VLBI)~\\cite{sudou}. In particular,  a newly found periodic flux variation suggests that this binary system will coalesce in 500 years~\\cite{iguchi}. Perhaps, we may conclude coalescence of binary black holes often takes place in the Universe.   However a direct evidence of black hole merger is still missing. One of the possibilities to {\\it see} the merger process is to observe shadows of black holes shone by the radiation from the accretion disc or star lights behind the black holes~\\cite{falcke}. Since two event horizons merge into one event horizon, we expect that the shadows must show very peculiar time evolution.   Observing these shadows, therefore,  should be compelling evidence of a coalescing black holes as well as provides an intriguing probe of general relativity with very strong gravity field.  As a first step toward the study of a realistic black hole binary, we have recently calculated  the shadows of the Kastor-Traschen (KT) \\cite{kt} cosmological multi-black hole solutions \\cite{nitta}. We have found that the shadows are deformed in the direction of the collision and that in addition to the shadow of each black hole, eyebrow-like structures appear as the black holes come close to each other. In this paper, we attempt to understand these structures analytically using the Majumdar-Papapetrou (MP) solution \\cite{mp}, the  static multi-black hole solution with the maximal charge to which the KT solution is reduced when the cosmological constant is zero.  The paper is organized as follows. In Sec. II, we present several analytic calculations of null geodesics in the extreme Reissner-Nordstr\\\"om and MP solutions in order to provide  analytic estimate of the eyebrow-like structures of the shadows as well as the deformation of the shadows in the MP solution. In Sec. III, extending our previous results \\cite{nitta}, we present numerical results of the shadows of the colliding/coalescing  black hole binary in the KT solution by changing the viewing angle of the distant observer and {compare them with the shadows of the MP solution.} Sec. IV summarizes the results. ",
        "conclusions": "We have studied the null geodesics in the static/dynamic multi-black hole solutions: Majumdar-Papapetrou (MP) solution and the Kastor-Traschen (KT) solution. We have calculated the shadows of these multi-black holes and found that the shadows have structures distinct from the mere superposition of the shadow of each black hole: the eyebrow-like structures outside the main shadows and the deformation of the shadows. We have presented analytic estimates of these structures using the MP solution to show that the width of the eyebrow is related with the distance between the black holes and that the shadows are deformed into ellipses due to the presence of the second black holes and the separation between the shadows is larger.   These analytic results help us to have qualitative understanding of the features of the shadows of colliding black holes which are studied in our previous paper. We expect that following two features of black hole shadows are general and appear in more realistic situation.  First one is the eyebrow-like structure which shows up during the merger process. Second is the deformation of the main shadow and the larger separation than the true distance.  These features in the shadows can be used as probes to find the multi-black hole system at the final stage of its merger process. For that purpose, we have presented the shadows of the colliding black holes in the KT solution by changing the direction of the observer to mimic the coalescence of the binary black holes. In order to study the shadows of a realistic black hole binary, the effects of the accreation disk should also be considered, which is left for our future study."
    },
    "1208/1208.3016_arXiv.txt": {
        "abstract": "We examine the metallicity distribution function (MDF) and fraction of carbon-enhanced metal-poor (CEMP) stars in a sample that includes 86 stars with [Fe/H]~$\\le$~$-$3.0, based on high-resolution, high-S/N spectroscopy, of which some 32 objects lie below [Fe/H]~=~$-$3.5. After accounting for the completeness function, the ``corrected'' MDF does not exhibit the sudden drop at [Fe/H] = $-$3.6 that was found in recent samples of dwarfs and giants from the Hamburg/ESO survey. Rather, the MDF decreases smoothly down to [Fe/H] = $-$4.1. Similar results are obtained from the ``raw'' MDF. We find the fraction of CEMP objects below [Fe/H]~=~$-$3.0 is 23 $\\pm$ 6\\% and 32 $\\pm$ 8\\% when adopting the \\citeauthor{beers05} and \\citeauthor{aoki07} CEMP definitions, respectively. The former value is in fair agreement with some previous measurements, which adopt the \\citeauthor{beers05} criterion. ",
        "introduction": "\\label{sec:intro}  Metal-poor stars provide critical information on the earliest phases of Galactic formation (see e.g., the reviews by \\citealt{beers05} and \\citealt{frebel11}). Their chemical abundances shed light upon the nature of the first stars to have formed in the Universe, and the nucleosynthesis which seeded all subsequent generations of stars.  This is the third paper in our series, which focuses upon the discovery of, and high-resolution, high signal-to-noise ratio (S/N) spectroscopic analysis of, the most metal-poor stars. Here we explore two key issues: the metallicity distribution function (MDF) and the fraction of carbon-enhanced metal-poor (CEMP)\\footnote{Initially defined as stars with [C/Fe] $\\ge$ +1.0 and [Fe/H] $\\le$ $-$2.0 \\citep{beers05}.} stars at lowest metallicities.  Any model purporting to explain the formation and evolution of our Galaxy must be able to reproduce the observed MDF.  The ingredients of such models include the initial mass function (IMF), nucleosynthetic yields, and inflow or outflow of gas.  Observations of the MDF can constrain these initial conditions and physical processes.  Since the early work by \\citet{hartwick76}, measurements of the MDF involve increasing numbers of stars with more accurate metallicity measurements (see e.g., \\citealt{laird88,ryan91}). One of the basic predictions of Hartwick's Simple Model of Galactic Chemical Enrichment is that the number of stars having abundance less than a given metallicity should decrease by a factor of ten for each factor of ten decrease in metallicity{\\footnote{While a number of chemical evolution models (e.g., \\citealt{kobayashi06}, \\citealt{karlsson06}, \\citealt{salvadori07}, \\citealt{prantzos08}, and \\citealt{cescutti10}) have improved upon the one-zone closed-box Hartwick model, the general behavior remains largely unchanged.}. \\citet{norris99} presented observational support for this suggestion, down to [Fe/H] $\\sim-4.0$, below which it appeared to be no longer valid. More recently, \\citet{schorck09} and \\citet{li10} presented MDFs of the Galactic halo using 1638 giant and 617 dwarf stars, respectively, from the Hamburg/ESO Survey (HES: \\citealt{hes}).  Below [Fe/H] = $-$2.5, the MDFs for dwarfs and giants were in excellent agreement. A prominent feature of both MDFs was the apparent lack of stars more metal-poor than [Fe/H] = $-$3.6. While a handful of such stars are known, the sharp cutoff in the MDF has important implications for the critical metallicity above which low-mass star formation is possible (e.g., \\citealt{salvadori07}).  More detailed studies of the MDF, and in particular the low-metallicity tail, are necessary to confirm and constrain the star formation modes of the first stars (e.g., \\citealt{bromm04}).  The HK survey \\citep{beers85,beers92} revealed that there is a large fraction of metal-poor stars with unusually strong CH $G$-bands indicating high C abundances. With the addition of numerous metal-poor stars found in the HES, the CEMP fraction at low metallicity has been confirmed and quantified, with estimates ranging from 9\\% \\citep{frebel06} to > 21\\% \\citep{lucatello06}.  These numbers are considerably larger than the fraction of C-rich objects at higher metallicity, the so-called CH and Ba stars, which account for only $\\sim$ 1\\% of the population.  The fraction is even larger at lowest metallicity: below [Fe/H] $<$ --4.5, 75\\% of the four known stars belong to the CEMP class \\citep{norris07,caffau11}.  To explain these large fractions, several studies argue that adjustments to the IMF are necessary (e.g., \\citealt{lucatello05,komiya07,izzard09}). \\citet{carollo12} offer an alternative interpretation for the increase of the CEMP fraction they observe in the range $-$3.0 $<$ [Fe/H] $<$ $-$1.5 in terms of a dependence of CEMP fraction on height above the Galactic plane. In their most metal-poor bin at [Fe/H] $\\sim$ --2.7, they report C-rich fractions of 20\\% and 30\\% for their inner- and outer-halo components, respectively (see their Figure 15). In their view, this can be accounted for by the presence of different carbon-production mechanisms (some not involving the presence of AGB nucleosynthesis) that have operated in the inner- and outer-halo populations.  An understanding of the CEMP stars is complicated by the fact that they do not form a homogeneous group: \\citet{beers05} define several CEMP subclasses (all of which have [C/Fe] $>$ +1.0) as follows: (i) CEMP-r -- [Eu/Fe] $>$ +1.0; (ii) CEMP-s -- [Ba/Fe] $>$ +1.0 and [Ba/Eu] $>$ + 0.5; (iii) CEMP r/s -- 0.0 $<$ [Ba/Eu] $<$ +0.5; and CEMP-no -- [Ba/Fe] $<$ 0.0.  \\citet{aoki10} shows that below [Fe/H] = --3.0, the CEMP stars are principally (90\\%) CEMP-no stars, while for [Fe/H] $>$ --3.0, the CEMP-s class predominates. These differences lie outside the scope of the present paper.  Here we seek to constrain only the fraction of CEMP stars at lowest abundance, [Fe/H] $<$ --3.0, and to compare the results with the fractions determined at higher abundances.  In Paper IV (Norris et al.\\ 2012b) we shall address the nature of the CEMP-no stars, which comprise the large majority of CEMP stars in our extremely metal-poor sample. ",
        "conclusions": ""
    },
    "1208/1208.4293_arXiv.txt": {
        "abstract": "Analysing all Galaxy and Mass Assembly (GAMA) galaxies within a factor two ($\\pm 0.3$ dex) of the stellar mass of the Milky Way (MW), there is a 11.9\\% chance that one of these galaxies will have a close companion (within a projected separation of 70 kpc and radial separation of 400 km/s) that is at least as massive as the Large Magellanic Cloud (LMC). Two close companions at least as massive as the Small Magellanic Cloud (SMC) are rare at the 3.4\\% level. Two full analogues to the MW-LMC-SMC system were found in GAMA (all galaxies late-type and star forming), suggesting such a combination of close together, late-type, star-forming galaxies is rare: only 0.4\\% of MW mass galaxies (in the range where we could observe both the LMC and SMC) have such a system. In summary, the MW-LMC-SMC system is a 2.7$\\sigma$ event (when recast into Gaussian statistics).  Using cross-correlation comparisons we find that there is a preference for SMC-LMC binary pair analogues to be located within 2 Mpc of a range of different luminosity groups. There is a particular preference is for such binaries to be located near LG luminosity systems. When these groups are subdivided into small magnitude gap and large magnitude gap subsets, the binaries prefer to be spatially associated with the small magnitude gap systems. These systems will be dynamically less evolved, but still offer the same amount of gravitational dark matter. This suggests that binaries such as the SMC-LMC might be transient systems, usually destroyed during vigorous merger events. Details of a particularly striking analogue to the MW-SMC-LMC and M31 complex are included. ",
        "introduction": "The Local Group (LG), and more specifically the Milky-Way (MW) is the most thoroughly explored dark matter complex in the Universe \\citep[e.g.][]{mate98,berg00,bens02,kara09,font11}. However, question marks remain over how typical the MW halo is in the context of the Universe and how unusual its galaxy occupation statistics are \\citep[e.g.][]{boyl11,love11,toll11,weis11}. It is important that we fully understand how representative the MW halo is since, by virtue of proximity, it will always be the environment that will contain the faintest known galaxies and the broadest range of galaxy masses. It will also be the halo from which we can derive the most information about its formation history. Knowing which satellites of the MW halo are typical within similar-mass similar-redshift haloes will either severely tighten or relax the predictive requirements of N-body semi-analytic galaxy formation codes. Currently it is acknowledged that simulations struggle to predict the full distribution of MW satellite galaxies, these problems are particularly manifest for the brightest satellites: Large and Small Magellanic Clouds (LMC and SMC) \\citep[e.g.][]{bens02,kopo09,okam10}.  We are set to learn a vast amount about the MW halo in the coming decades. In the near future GAIA \\citep{wilk05} will measure space motions and properties for 2 billion stars in the LG which includes all known member galaxies. Amongst likely discoveries, we will learn about dynamical equilibrium, or lack of it, for the first time. Building up to these hugely detailed surveys it is important we discover where the MW halo fits into the bigger picture. Only then can we apply what we know about the MW to the $\\Lambda$CDM (or some variant) model of the Universe. Combining near-field cosmology (LG scale) and far-field cosmology (redshift surveys) is key to completing the full picture of galaxy formation \\citep{free02}.  This work puts the investigation of the MW halo into an observational cosmological context by using data from the Galaxy and Mass Assembly project (GAMA). GAMA is a multi-wavelength photometric and spectroscopic survey, and was designed to answer questions about how matter has assembled on a huge variety of scales: filaments, clusters, groups and galaxies. The first phase of the redshift survey was conducted on the AAT (known as GAMA-I) and these data are used in this work \\citep{driv11}. In this work we use GAMA redshifts to search for close companions to MW mass galaxies. These systems will be Milky-Way Magellanic Cloud Analogues (MMAs from here). We use this sample to construct statistics on the rarity of SMC and LMC type (star forming late-type galaxies) close companions around $L^*$ late-type moderate star-formation rate spiral galaxies like our own MW.  In Section \\ref{sec:Data} we discuss the data used in this work in detail. In Section \\ref{sec:LGA} we present the statistics for finding MW-LMC and MW-LMC-SMC type systems, allowing us to quantify the apparent rarity of MW like systems. In Section \\ref{sec:Cross} we investigate the environment that SMC-LMC type binaries are most commonly located in, and relate this to some of the defining characteristics of the LG.  Data for the LG was calculated using distance indicators without any $H_0$ dependence. As such it is appropriate (and consistent with the main body of LG literature) to convert GAMA data into true distance. To make the appropriate conversions we take the latest WMAP 7 cosmology: $\\Omega_M=0.27$, $\\Omega_\\Lambda=0.73$ and $H_{0}=70\\,{\\rm km}\\,{\\rm s}^{-1}$ \\citep{koma11}.  \\vspace{-8mm} ",
        "conclusions": "The major findings of this work are summarised below:  \\begin{itemize}   \\item Analysing all galaxies within 0.3 dex of the stellar mass of the MW, there is a 11.9\\% (11.2\\%--12.8\\%) chance that it will have a close companion (within a projected separation of 70 kpc and radial separation of 400 km/s) that is at least as massive as the LMC. This is consistent with analyses by \\citet{boyl11} of the Millenium II simulation and by \\citet{jame11} of H$\\alpha$ imaging around luminous spiral galaxies.  \\item Limiting the sample to those galaxies where the SMC would be observable we find 3.4\\% (2.7\\%--4.5\\%) of galaxies have two companions at least as massive as the SMC.  \\item Only two full analogues to the MW-LMC-SMC system were found in GAMA, suggesting such a combination of late-type, close star-forming galaxies is quite rare: in GAMA only 0.4\\% (0.3\\%--1.1\\%) of MW mass galaxies have such a system (a 2.7$\\sigma$ event). In terms of space density, we find $1.1 \\times 10^{-5}$Mpc$^{-3}$ full analogues in GAMA (in a volume of $1.8 \\times 10^{5}$Mpc$^3$). The best example found shares many qualitative characteristics with the MW system. The brightest pair galaxy has spiral features, as does the bigger minor companion. The minor companions are $\\sim$40 kpc in projected separation, so not in a close binary formation like the SMC and LMC.  \\item Selecting systems that are close binaries like the SMC-LMC pair (MCAs), we find that they are preferentially located in close proximity (or within) systems that have a similar total flux to the LG ($\\Sigma M_r = -22.5 \\pm 0.5$ mag).  \\item Subdividing the preferential group type into those with large and small magnitude gaps, we find MCAs are more spatially associated with groups that have a small magnitude gap. This suggests that a quiet recent merger history improves the likelihood of the Magellanic Clouds being visible in the LG. The best MMA analogue found in GAMA also has a close $L^*$ spiral companion galaxy.  \\end{itemize}      \\vspace{-8mm}"
    },
    "1208/1208.6296_arXiv.txt": {
        "abstract": "In order to fully understand the gravitational collapse of molecular clouds, the star formation process and the evolution of circumstellar disks, these phenomena must be studied in different Galactic environments with a range of stellar contents and positions in the Galaxy. The young massive association Cygnus~OB2, in the Cygnus-X region, is an unique target to study how star formation and the evolution of circumstellar disks proceed in the presence of a large number of massive stars. We present a catalog obtained with recent optical observations in $r,\\,i,\\,z$ filters with OSIRIS, mounted on the $10.4\\,m$ GTC telescope, which is the deepest optical catalog of Cyg~OB2 to date. \\par The catalog consist of 64157 sources down to $M=0.15\\,M_{\\odot}$ at the adopted distance and age of Cyg~OB2. A total of 38300 sources have good photometry in all three bands. We combined the optical catalog with existing X-ray data of this region, in order to define the cluster locus in the optical diagrams. The cluster locus in the $r-i$ vs. $i-z$ diagram is compatible with an extinction of the optically selected cluster members in the $2.64^m<A_V<5.57^m$ range. We derive an extinction map of the region, finding a median value of $A_V=4.33^m$ in the center of the association, decreasing toward the north-west. In the color-magnitude diagrams, the shape of the distribution of main sequence stars is compatible with the presence of an obscuring cloud in the foreground at $\\sim850\\pm 25\\,pc$ from the Sun. ",
        "introduction": "In recent years our knowledge of the star formation process and the early phase of stellar evolution has increased remarkably.  This progress has been driven by a large number of deep spectroscopic and photometric observations at high spatial resolution of star forming regions in our Galaxy and in nearby galaxies.  High performance telescopes, such as the NASA Hubble and Spitzer Space Telescopes, and more recently the ESA Herschel Space Observatory, allow us to study the formation of low- and high-mass stars, the evolution of circumstellar disks, and the collapse of protostellar cores in unprecedented detail. X-ray observations with space telescopes such as {\\it Chandra} and XMM-Newton have also been crucial, unveiling the populations of star forming regions down to sub-solar masses, and probing the high energy processes in young stars. \\par Despite this progress, few single star forming regions in our Galaxy have provided the opportunity to study star and planet formation over a large range of stellar masses and in presence of very massive stars, which affect the star formation process and the evolution of their parental cloud mostly thank to their intense ionizing flux. The study of such massive young clusters, in fact, is hampered by their large distance from the Sun, with only a few moderately massive clusters closer than $1.5\\,kpc$.  \\par One exception that has only fairly recently been recognized is the Cygnus~OB2 association \\citep{Red67}, in the Cygnus-X giant molecular cloud.  It is the richest and most massive OB2 association within $2\\,kpc$ of the Sun. It harbors a large number of massive stars and an extensive young population of low-mass pre-main sequence stars, likely formed in different episodes. The first spectroscopic study devoted to the massive members of this region \\citep{Red67} identified about 300 OB members, and in subsequent studies this estimate has increased. In a 2MASS study, \\citet{Kno00} used statistical arguments to suggest the presence of 2600 OB stars and 120 O stars and to estimate a total mass of $\\left(4-10 \\right) \\times 10^4\\, M_{\\odot}$.  This mass led the author to the conclusion that Cyg~OB2 is a young globular cluster in the Milky Way. A near-infrared spectroscopic survey by \\citet{Come02} found slightly fewer O stars but confirmed this conclusion.  While \\citet[][see also \\citealt{Dre08}]{Han03} finds a somewhat lower total mass. The enormous scale of Cyg~OB2 is exemplified by it containing two of the few O3 stars known in our Galaxy \\citep{Wal73}, together with B supergiant stars \\citep{Mas91,Negu08}.  Other massive objects, such as Wolf-Rayet stars \\citep{Nie98} and the Luminous Blue Variable G79.29+0.46 \\citep{Hig94} are found in the Cygnus-X complex in the proximity of the association. \\par The first studies of the intermediate- and low-mass population of the association counted several thousands of candidate members \\citep{Red67,Kno00}, but most recent and reliable estimates based on {\\it Chandra} X-ray observations of the central cluster indicate a population of 1000-1500 low-mass members down to subsolar masses \\citep{Alba07,Wri09}, among which there are several strong $H\\alpha$ emitting objects \\citep{Vink08}. \\par The first estimate of the distance to Cyg~OB2 was $2.1\\,kpc$ \\citep{Red67}. \\citet{Mas91} found $1.8\\,kpc$ based on a combined photometric and spectroscopic study, while the most recent estimate based on spectroscopy of the most massive members of Cyg~OB2 is $1.45\\,kpc$ \\citep{Han03}. \\citet{Rygl11} reported a distance of the Cygnus-X complex of $1400\\,pc$, derived from the trigonometric parallaxes of five star forming region in the complex. We adopt a distance of $1450\\,pc$ in this work.  Cyg~OB2 then lies behind the nebulosity associated with the Great Cygnus Rift and, despite its relative proximity to the Sun, the association is significantly absorbed. \\par A picture has emerged in the last few years of different episodes of star formation in and around the main center of Cyg~OB2. \\citet{Han03} noted the presence of both massive stars younger than $\\sim2\\,Myr$, and of B stars that appeared older for the adopted association distance, or would need to be in the foreground. \\citet{Dre08} identified a population of A-type stars south of the main association with an age of $5-7\\,Myr$. Since then, \\citet{Wri00} have shown that the NIR color-magnitude diagram of the region indicates both a $\\sim3\\,Myr$ and a $\\sim5\\,Myr$ old population of stars down to stellar masses of $\\sim1\\,M_{\\odot}$, while \\citet{Wright12} discovered a population of candidate proplyds\\footnote{Stars with protoplanetary disks surrounded by a photoevaporating envelope of gas \\citep{Dell94}} southward Cyg~OB2 suggesting that new episodes of star formation occurred. \\par In this paper we present the deepest and highest spatial resolution optical observations of Cyg~OB2 to date. In sections \\ref{inst_sect} and \\ref{phot_sect} we describe the instrument setup, the observations, the data reduction, and the photometric calibration in the AB system; in section \\ref{cat_sect} we describe the final catalog, consisting of 64157 optical sources, among which 38300 have good photometry in all three bands. In section \\ref{match_sec} we cross-correlate the OSIRIS catalog with the IPHAS data and with the newly released SDSS DR8 data of this region, and with the existing {\\it Chandra}/ACIS-I catalog of Cyg~OB2. Finally, in sections \\ref{av_sec} and \\ref{cm_sec} we present the color-color and color-magnitude diagrams and we use them to study the extinction and age of association members and the distance of the foreground nebulosity. In a forthcoming paper, the optical catalog described in this paper will be combined with new deep X-ray observations of Cyg~OB2 (the $1.08\\,Msec$ {\\it Chandra} Cygnus~OB2 Legacy Survey, P.I. J.~Drake\\footnote{http://www.cygob2.org/}), to obtain the deepest list of stars associated with this unique star forming region. \\par ",
        "conclusions": "We have analyzed new GTC/OSIRIS optical observations in $r^{\\prime}i^{\\prime}z^{\\prime}$ bands of a field of size $41^{\\prime}\\times 41^{\\prime}$ approximately centered on Cyg~OB2. The resulting catalog contains the photometry of 64157 optical sources, among which 38300 have good photometry in all three bands. The catalog reaches $r^{\\prime}=25^m$, 5 magnitudes deeper than the existing IPHAS catalog and 3 than the SDSS DR8 data. This limit corresponds to a $3.5\\,Myr$ star with $M=0.15\\,M_{\\odot}$ at the adopted distance and extinction of Cyg~OB2. The stellar density of stars with good photometry varies significantly from north southward, with the maximum density corresponding to a cavity in the cloud in the north-west part of the observed area. \\par We cross-correlated the OSIRIS catalog with an existing {\\it Chandra} ACIS-I catalog, of the central and north-west area of the field, in order to define the cluster locus in the diagrams. This yielded 1407 optical sources with X-ray counterparts, almost twice the number of X-ray sources with optical counterparts identified in previous works in these fields. Using the stars with detections in both the OSIRIS and IPHAS catalogs, we derived a suitable color transformation between the IPHAS and OSIRIS photometric systems (Vega and AB, respectively). This transformation converts properly the colors of a G2V star from the AB to the Vega system. \\par The $r-i$ vs. $i-z$ color-color diagram shows a clear cluster sequence which can be encompassed by a $3.5\\,Myr$ isochrone with extinction between $2^m$ and $6^m$.  This extinction turns out to be the main range of extinction of cluster members.  We also found that the mean extinction decreases from the central cluster northward: in the central field we found a median extinction $A_V=4.33^m$, with a large differential reddening across the field; while in the north-west {\\it Chandra} field we found a median extinction $A_V=3.21^m$, with a more peaked extinction distribution.  We derived an extinction map of the region that reveals evidence that the most obscured regions are also those with the largest density of low-mass and massive member stars. \\par The color-color diagram also exhibits a gap between the foreground objects and the cluster locus that is induced by a steep increase of the visual extinction. This is due to the presence of a dense nebulosity along this line of sight. The fit of the zero age main sequence with the foreground population in the $g$ vs. $g-i$ diagram suggests a distance of $850\\pm25\\,pc$ for this cloud, which turns out to be in the foreground and not associated with Cyg~OB2."
    },
    "1208/1208.1680_arXiv.txt": {
        "abstract": "We present the relation between the ($z-$ and $k-$corrected) spectral lags, $\\tau$, for the standard \\textit{Swift} energy bands $50-100$ keV and $100-200$ keV and the peak isotropic luminosity, $L_{\\mathrm{iso}}$ (a relation reported first by Norris et al.), for a subset of $12$ long \\textit{Swift} GRBs taken from a recent study of this relation by Ukwatta et al. The chosen GRBs are also a subset of the Dainotti et al. sample, a set of {\\em Swift} GRBs of known redshift, employed in establishing a relation between the (GRB frame) luminosity, $L_X$, of the shallow (or constant) flux portion of the typical XRT GRB-afterglow light curve and the (GRB frame) time of transition to the normal decay rate, $T_{\\mathrm{brk}}$. We also {present} the $L_X-T_{\\mathrm{brk}}$ relation using only the bursts common in the two samples. The two relations exhibit a significant degree of correlation ($\\rho = -0.65$ for the $L_{\\mathrm{iso}}-\\tau$ and $\\rho = -0.88$ for the $L_{X} - T_{\\mathrm{brk}}$ relation) and have surprisingly similar best-fit power law indices ($-1.19 \\pm 0.17$ for $L_{\\mathrm{iso}}-\\tau$ and $-1.10 \\pm 0.03$ for $L_{X} - T_{\\mathrm{brk}}$). Even more surprisingly, {we noted that although $\\tau$ and $T_{\\mathrm{brk}}$ {represent} different GRB time variables}, it appears that the first relation ($L_{\\mathrm{iso}}-\\tau$) extrapolates into the second one for {timescales} $\\tau \\simeq T_{\\mathrm{brk}}$. This fact suggests that these two relations have a common origin, which we conjecture to be kinematic. This relation adds to the recently discovered relations between properties of the prompt and afterglow GRB phases, indicating a much more intimate relation between these two phases than hitherto considered. ",
        "introduction": "Gamma-ray Bursts (GRBs) are extremely bright explosions, with isotropic luminosities exceeding $\\sim 10^{54}$ erg/sec, durations in the range $\\sim 0.1 - 1000$ sec and energy of peak luminosity in the $\\gamma-$ray regime, $E_p \\sim 1$ MeV, hence their name. They are believed to originate in the collapse of stellar cores or the mergers of neutron stars, processes which result in jet-like relativistic outflows of Lorentz factors $\\Gamma \\sim 300$, whose kinetic energy is converted efficiently into radiation at distances $r \\sim 10^{15} - 10^{18}$ cm at a relativistic blast wave (RBW) to produce the observed events (for a review see Piran 2005). Following their most luminous, prompt, $\\gamma-$ray emission, they shift into the afterglow phase with peak luminosity in the X-ray band and duration of $\\sim 10^5$ s, in which their localization can be refined and optical detection can provide their redshift.  The theory of  RBW  slow-down indicated a smooth power law decay $\\propto t^{-1}$ for the flux of their afterglow X-ray light curves, and indeed the pre-{\\em Swift} sparsely sampled ones appeared consistent with such a behavior. However, their more densely sampled X-ray light curves with the XRT aboard {\\em Swift} uncovered significant deviations from this behavior. So, following the prompt emission in $\\gamma-$rays, the typical XRT afterglow consists \\citep{nousek} of a much steeper flux decline ($\\propto t^{-3}\\ \\mbox{to}\\ t^{-6}$), followed by a  $10^2 - 10^5$ sec period of nearly constant flux, followed finally at $t = T_{brk}$ by the more conventional power-law decline $\\simeq t^{-1}$. In addition, {\\em Swift} follow-ups discovered also occasional flares on top of these light curves, as late as $\\sim 10^5$ sec since the Burst Alert Telescope (BAT) trigger.  In their prompt phase GRBs exhibit a broad light curve diversity and a large variance in their (estimated) RBW Lorentz factors ($\\Gamma \\sim 100-1000$). These properties, along with the non-thermal character of their spectra, suggest that at least in this phase, GRBs are not likely to provide well defined, underlying systematics that would allow a probe of the physics of prompt emission. However, a number of correlations has been found between observables of the prompt phase \\citep{rie01,sch03,ghi06,ama09}, whose origin remains still largely unaccounted for.   One of the first such correlations, made possible only after the determination of the GRB redshifts by their afterglow emission, is that between the burst peak isotropic luminosity, $L_{\\rm iso}$, and the spectral lag, $\\tau$, between different energy bands in the GRB spectrum. This relation has been studied in detail by a number of authors \\citep{nor00,nor02,geh06,sch07,sta08,hak08,ukw10,ukw11,wan11} for data sets obtained by different instruments aboard different missions.  The general conclusion of all these studies has been an anti-correlation $L_{\\mathrm{iso}} \\propto \\tau^{-a}$ with a value of $a \\sim 0.77 - 1.8$. The spectral lag is defined as the difference in the arrival time of high and low energy photons, and is taken to be positive when the time of arrival of high energy photons precedes that of low energy photons.  Normally the spectral lag is extracted between two arbitrary energy bands in the observer frame and then corrected for the time dilatation effect ($z-$correction) by multiplying the lag value in the observer frame by $(1+z)^{-1}$. Moreover the observed energy bands correspond to different energy bands at the GRB source frame for different redshifts, and so one needs to take into account this energy dependent factor ($k-$correction). Using the assumption that the spectral lag is proportional to the pulse width which in turn is proportional to the energy, \\citet{geh06} approximately corrected for this effect by multiplying the lag value in the observer frame by $(1+z)^{0.33}$. Alternatively the $k-$correction can be done by extracting the spectral lags in the GRB source frame. This is accomplished by choosing two energy bands in the source frame and then projecting these in the observer frame using the relation $E_{\\mathrm{obs}} = E_{\\mathrm{source}}/(1 + z)$, such that the projected energy bands lie in the \\textit{Swift}  BAT energy range $15-350$ keV. This alternative method of extracting the spectral lags in the source frame rather than the observer frame, was recently used by \\citet{ukw11} to obtain a similar result for the $L_{\\mathrm{iso}}- \\tau$ relation, but with a higher degree of correlation between the variables.  The physical origin of the lag-luminosity relation is still unclear. Some, including \\citet{sal00} and  \\citet{iok01}, attributed this to a kinematic effect; \\citet{sch04}, following an in depth analysis of existing proposals, concludes in favor of the time evolution of the emitting electrons, on the basis of a correlation between the energy of the GRB peak emission and the burst instantaneous flux. However, irrespective of its physical origin, this relation is a useful tool in GRB science, not only for its use in distinguishing between long and short bursts [with long bursts {in general} exhibiting larger lags than short ones; {although it has been shown \\citep{geh06,hak07} that this may not be necessarily true for every long GRB}], but also for its implementation, together with other relations between GRB variables, in extending the Hubble diagram to higher redshifts \\citep{sch07,wan11}.  An altogether different correlation, pertaining to the GRB afterglow phase, has been reported recently by \\citet{dai10}. This work presents a correlation between the X-ray luminosity, $L_X$, of the plateau (or shallow decay) phase and the source frame break time $T_{\\rm brk}$ in the XRT light curves of long GRBs. Using a sample of 62 \\textit{Swift} long GRBs a correlation of the form $\\log L_{X} = \\log a + b \\log T_{\\mathrm{brk}}$, with $\\log a = 51.06\\pm 1.02$ and $b = -1.06^{+0.27}_{-0.28}$ was obtained. A similar but steeper correlation $(b = -1.72^{+0.22}_{-0.21})$ was also obtained for a small group of GRBs which belong to the intermediate class (IC) \\citep{nor06} between short and long ones, indicating that these may behave in a different way than the long GRBs. There have been claims \\citep{can11} that the Dainotti relation is just a selection effect due to the flux detection limit for \\textit{Swift}'s XRT which prevents clear observation of faint light curves from high redshift GRBs. This possibility was later investigated by \\citet{dai11a} who showed that there is no systematic bias against faint plateaus at high $z$, thus confirming the existence of this relation. Moreover \\citet{dai11b} have obtained a number of significant correlations between the afterglow phase X-ray luminosity parameter $L_{X}$ and prompt emission parameters such as the isotropic energy $E_{\\mathrm{iso}}$, peak energy $E_{\\mathrm{peak}}$ and the variability parameter $V$ \\citep{nor00}.\\\\  In this work we use a sample of 14 GRBs which are common in the \\citet{ukw10} and \\citet{dai10} studies, to obtain and compare their lag-luminosity and break time - X-ray luminosity relations in the GRB source frame after doing the necessary $k$- and $z$- corrections. The structure of the paper is as follows. In Section 2 we discuss briefly the correlations and computational procedures involved in the spectral lag - isotropic luminosity relation of \\citet{ukw10} together with the X-ray luminosity of the GRB shallow afterglow phase and its break time of \\citet{dai10}. In Section 3 we present our results and then in Section 4 we summarize our findings and conclusions. ",
        "conclusions": "In the sections above we reviewed correlations between the luminosities and time scales of two different stages in the GRB development, namely the prompt emission and the shallow decay stage of their afterglow. We then reproduced the correlations between $L - \\tau$ for the prompt emission and $L_X - T_{\\rm brk}$ for the afterglow, for the GRBs common in the data sets used by \\citet{ukw10} and \\citet{dai10}, from the data {already} present in the literature. {We have shown that although our GRB sample is small, both relations are consistent (in terms of power-law index and correlation coefficient) with the previous relations obtained using larger samples. Moreover we have shown that these two relations extrapolate very well into each other and give a much tighter relation (Eq. (9)) than the individual relations obtained so far.}  For the first time we also noted that although the relations in Equations (\\ref{LT}) and (\\ref{LxTa}) represent different stages in the GRB evolution, their power-law indices are surprisingly similar. Yet from our data the source frame corrected spectral time lag and break time $T_{\\mathrm{brk}}$ do not appear to be correlated. The fact that their normalizations are such that they extrapolate into each other, suggests that prompt and afterglow properties are interrelated. This fact could have been actually surmised from the original lag treatment of \\citet{nor00} and the results of \\citet{dai10}. The reader can easily confirm that the relation of \\citet{nor00} extrapolates into that of \\citet{dai10}. A correlation between the prompt and afterglow phases has also been explored by \\citet{sal02} who obtained a correlation between the spectral lag $\\tau$ and rest frame jet-break time $\\tau_{j}$ given by \\begin{equation} \\tau_{j} = 28^{+18}_{-11}\\left(\\frac{\\tau}{1\\mathrm{s}}\\right)^{0.89 \\pm 0.12}\\ \\mathrm{days}, \\label{salmonson} \\end{equation} using a sample of seven BATSE gamma-ray bursts. Another correlation between prompt and afterglow quantities was recently obtained by \\citet{mar12}, who obtained a three-parameter correlation between the rest frame isotropic energy in the prompt phase $E_{\\gamma,\\mathrm{iso}}$, the peak of the prompt emission energy spectrum $E_{\\mathrm{pk}}$, and the X-ray energy emitted in the $0.3 - 10$ keV observed energy band $E_{X,\\mathrm{iso}}$ given by \\begin{equation} E_{X,\\mathrm{iso}} \\sim \\frac{E_{\\gamma, \\mathrm{iso}}}{E_{\\mathrm{pk}}^{3/4}}. \\end{equation} It was shown that this relation is robust and independent of the definition of $E_{X,\\mathrm{iso}}$. Moreover \\citet{mar12} and \\citet{ber12} showed that this three-parameter relation is shared by both long and short gamma ray bursts and {also} claim that the physical origin of such a relation is related to the outflow Lorentz factor.  At this stage we do not intend to speculate on the possible physics underlying the correlations of (\\ref{LT}) and (\\ref{LxTa}). Instead, we present a brief review of proposed explanations found in the literature. Then we conclude that, if indeed the underlying physics is common, as Figure \\ref{fig3} suggests, the apparently common origin of the two effects is basically kinematic.  In the case of the spectral lag-luminosity relation one possible explanation for the observed lags involves the spectral evolution during the prompt phase \\citep{der98,koc03,ryd05} in which due to cooling effects $E_{\\mathrm{pk}}$ shifts towards a lower energy band so that the temporal peak of the corresponding light curve will also shift to lower energies, thereby resulting in the observed lag. Another explanation for the spectral lag-luminosity relation is based purely on kinematic effects \\citep{sal00,sal02,iok01,der04,she05,lu06}, where the peak luminosity $L_{\\mathrm{pk}}$ and spectral lag $\\tau$ depend on a single kinematic variable \\begin{equation} D = \\frac{1}{\\Gamma(1 - \\beta\\cos\\theta)(1 + z)}, \\label{doppler} \\end{equation} which represents the Doppler factor for ejecta in a jet, moving at an angle $\\theta$ from the line of sight with velocity $\\beta \\equiv v/c$ at redshift $z$. This kinematic variable relates a proper timescale $\\tau$ in the GRB rest frame to the observed timescale $t$ given by \\begin{equation} t = \\frac{\\tau}{D}, \\label{timescales} \\end{equation} so that if the spectral lag is due to some decay time scale $\\Delta\\tau$ in the GRB restframe, then this will become $\\Delta t = \\Delta\\tau/D$ in the lab frame. Moreover assuming a power-law spectrum with a low end of the form $\\phi(E) \\propto E^{-\\alpha}$, where $\\alpha$ is the low energy spectral index, \\citet{sal00} showed that the peak luminosity varies as \\begin{equation} L_{\\mathrm{pk}} \\propto D^{\\alpha}, \\label{lpk} \\end{equation} with $\\alpha \\approx 1$ \\citep{pre98}, so that equations (\\ref{timescales}) and (\\ref{lpk}) lead to the lag luminosity relation. The same argument was used by \\citet{sal02} to explain their correlation in equation (\\ref{salmonson}), where in this case the jet-break time $\\tau_{j} \\propto 1/D$.  In this approach, the dependence of luminosities and observed timescales on the single variable $D$ leads to the conclusion that the observed variety among GRBs has a kinematic origin, brought through variation of the viewing angle $\\theta$ or the Lorentz factor $\\Gamma$ profile of the jet $\\Gamma$, or both. So for example, \\citet{sal00} showed that the lag-luminosity relation is due only to a variation in the line-of-sight $\\Gamma$ among bursts, with high $\\Gamma$ bursts having smaller spectral lags and low $\\Gamma$ bursts exhibiting longer ones.  On the other hand \\citet{iok01} showed that the lag-luminosity relation can be explained by variation in the observer angle, $\\theta_{v}$, from the axis of the jet, using a simple jet in which $\\Gamma = \\mathrm{const.}$ for $\\theta < \\theta_{j}$, and zero emission for $\\theta > \\theta_{j}$, where $\\theta_{j}$ is the opening angle of the jet. In this case the lags arise due to the path difference between the near and far edges of the emitting region such that bright (dim) bursts with short (long) spectral lags correspond to small (large) viewing angle.  An explanation for the anticorrelation between the duration of the intrinsic plateau phase of the GRB light curve and X-Ray luminosity has been proposed by \\citet{dal11} using a model in which energy from a long-lived central engine is continuously injected to balance the radiative losses. These radiative losses will be stronger for higher luminosity, thus leading to shorter plateaus. Another explanation which is  based on the kinematic effect discussed above was proposed by \\citet{eic06}, who claimed that the flat (or sometimes slightly rising) decay phase of the afterglow lightcurve results from the combination of the decaying tail of the prompt emission and early afterglow observed at viewing angles slightly outside the edge of the jet. For such ``offset'' viewing angles the afterglow flux initially rises at early times when the beaming of radiation away from the line of sight gradually decreases, then rounds off as the beaming cone expands to include the line of sight, and finally joins the familiar decaying light curve.  Clearly, the relations given by equations (\\ref{LT}), (\\ref{LxTa}) and (\\ref{LtotT}) call for further analysis with larger data sets to determine whether the indices and normalization of these relations are indeed consistent with those presented above. Since the relations {in} (\\ref{LT}), (\\ref{LxTa}) correspond to {the} prompt and afterglow phases of the GRB evolution, the similarity of their power-law indices and normalizations (they extrapolate into each other in Figure (\\ref{fig3})) is an indication that a common process, probably kinematic, is responsible for the observed spectral lags and the shallow decay phase of the afterglow light curve. As discussed above, both these relations were attributed individually \\citep{iok01,eic06} to the same kinematic process, namely viewing the grb jets at ``off-beam'' lines of sight. The results presented in this paper are in accordance with this explanation."
    },
    "1208/1208.5776_arXiv.txt": {
        "abstract": "We report on a heterodyne receiver designed to observe the astrophysically important neutral atomic oxygen [OI] line at 4.7448 THz.  The local oscillator is a third-order distributed feedback Quantum Cascade Laser operating in continuous wave mode at 4.741 THz. A quasi-optical, superconducting NbN hot electron bolometer is used as the mixer.  We recorded a double sideband receiver noise temperature (T$^{DSB}_{rec}$) of 815 K, which is $\\sim$7 times the quantum noise limit ($\\frac{\\rm{h}\\nu}{2\\rm{k_B}}$) and an Allan variance time of 15 s at an effective noise fluctuation bandwidth of 18 MHz.  Heterodyne performance was confirmed by measuring a methanol line spectrum. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0479_arXiv.txt": {
        "abstract": "{The recent realization that most stars form in clusters, immediately raises the question of whether star and planet formation are influenced by the cluster environment. The stellar density in the most prevalent clusters is the key factor here. Whether dominant modes of clustered star formation exist is a fundamental question. Using near-neighbour searches in young clusters Bressert et al. (2010) claim this not to be  the case. They conclude that - at least in the solar neighbourhood - star formation is continuous from isolated to densely clustered and that the environment plays a minor role in star and planet formation.} {We investigate under which conditions near-neighbour searches in young clusters can distinguish between different modes of clustered star formation. } {Model star clusters with different memberships and density distributions are set up and near-neighbour searches are performed. We investigate the influence of the combination of different cluster modes, observational biases, and types of diagnostic on the results. } {We find that the specific cluster density profile, the relative sample sizes, limitations in observations and the choice of diagnostic method decides whether modelled modes of clustered star formation are detected by near-neighbour searches.  For density distributions that are centrally concentrated  but span a wide density range (for example, King profiles) separate cluster modes are only detectable under ideal conditions (sample selection, completeness) if the mean density of the individual clusters differs by at least a factor of $\\sim$65. Introducing a central cut-off can lead to underestimating the mean density by more than a factor of ten especially in high density regions. Similarly, the environmental effect on star and planet formation is underestimated for half of the population in dense systems.} {Local surface density distributions are a very useful tool for single cluster analysis, but only for high-resolution data. However, a simultaneous analysis of a sample of cluster environments involves effects of superposition that suppress characteristic features very efficiently and thus promotes erroneous conclusions. While multiple peaks in the distribution of the local surface density in star forming regions imply the existence of different modes of star formation, the reverse conclusion is {\\em not} possible. Equally, a smooth distribution is {\\em not} a proof of continuous star formation, because such a shape can easily hide modes of clustered star formation.} ",
        "introduction": "Most stars form in proximity to other stars within embedded clusters rather than being uniformly distributed throughout molecular clouds \\citep{1999A&A...342..515T,2000AJ....120.3139C,2003ARA&A..41...57L,2003AJ....126.1916P,2007prpl.conf..361A}. The density in young clusters in the Milky Way varies over many orders of magnitude from $\\ll$ 1 stars/pc$^3$ in relatively sparse clusters to $>$10$^5$ stars/pc$^3$ in the central areas of dense clusters.  The key factor in determining the relative importance of the environment for star and planet formation is the stellar density in the young clusters. Stars forming in the sparse cluster environments are largely unaffected by the presence of their fellow cluster members. By contrast, one can expect a strong influence on star and planet formation by the environment in the densest of these young clusters. Theoretical investigations predict that this environmental influence on star formation might manifest itself in a different initial mass function \\citep{2006MNRAS.368..141F,2006ApJ...652L.129P,2012MNRAS.tmp.2706M}, the binary fraction \\citep{2011MNRAS.417.1684M,2011A&A...528A.144K} and the disc frequency in high stellar density environments \\citep{2001MNRAS.325..449S,2005A&A...437..967P,2006A&A...454..811P,2006ApJ...642.1140O,2010A&A...509A..63O}. Observations have found indications of a dependence of these properties on the stellar density in young clusters \\citep{1998ApJ...492..540H,2008ApJ...675.1319H,2010ApJ...718..810S}. In dense clusters interactions might lead to a lower disc frequency  resulting in a lower planetary system rate and different properties in the planetary system.  For the stellar population as a whole the question is whether the properties of prestellar cores largely determine the stellar properties as in isolated star formation \\citep{2004ApJ...601..930S,2005MNRAS.359..211L,2006ApJ...641L.121T} or whether most stars form in a more dynamic way, where external forces and interactions dominate over initial conditions \\citep[e.g.][]{2001MNRAS.323..785B,2006MNRAS.370..488B}.  So a fundamental question of current star formation research is whether there exists a type of cluster (stellar emembership, density) that is the dominant environment for star formation? At first sight it would seem easy enough to answer this by simply collecting cluster data and determine the distribution of the mean density in young clusters. However, this is hindered by a number of obstacles. Most star formation occurs inside the spiral arms and close to the center of the Milky Way where it is difficult to identify clusters due to our position  within the plane of the Galactic disc. This means we have nothing like a complete census of the young clusters in the Milky Way. In principle, looking at nearby galaxies should help, but the larger distance means that the detection of low-mass clusters is hindered by their low luminosity.  There are different strategies for tackling the issue indirectly. One way is to look at the initial mass function \\citep[e.g.][]{2007prpl.conf..149B,2008ApJ...681..771M,2009MNRAS.392.1363B,2012ApJ...748...14D} or the binary development \\citep[e.g.][]{1994A&A...286...84D,1998ApJ...499L..79B,1999A&A...341..547D,2008AJ....135.2526C,2009ApJ...707.1533F,2011A&A...528A.144K,2011MNRAS.417.1684M} in different types of young clusters and compare them to the field properties. Similiarities i% are then interpreted as signs for a dominant cluster mode. However, since many cluster modes contribute simultaneously, a one-to-one relation is difficult to establish.  Another method is to measure the local surface density distribution in different cluster environments. Recently several observational studies \\citep[e.g.][]{2009ApJS..184...18G,2010MNRAS.409L..54B,2012ApJ...745..131K} tried to answer above questions by analyzing large samples of young stellar objects concerning their local surface density, $\\Sigma$, predominantly in the solar neighbourhood. Here it is argued that if different discrete modes existed they should manifest themselves as peaks in a surface density distribution \\citep[e.g.][]{1993ApJ...412..233S,2000AJ....120.3139C,2004MNRAS.350.1503W,2009ApJ...696...47W,2010MNRAS.409L..54B}.  This simple approach has the advantage that it does not rely on the definition of stellar groups, but uses the local separation from the star to its nearest neighbours. The local surface density is simply defined as $ \\Sigma = (n-1)/(\\pi r_n^2)$ where $n$ is the considered number of nearest neighbours including the star itself and $r_n$ is the distance to the $n$-th neighbour. Higher values of $n$ give a lower spatial resolution, but smaller fractional uncertainty \\citep{1985ApJ...298...80C,2009ApJS..184...18G}.  Using this method \\citet{2010MNRAS.409L..54B} found no peaks in the combined surface density distribution of several clusters in the solar neighbourhood (see their Fig.~1). They concluded from the absence of such peaks that star formation is continuous from isolated to densely clustered. In addition, they deduce a mean stellar surface density of 20~stars/pc$^2$ for the star forming regions in the solar neighbourhood and concluded that the environment plays a minor role in star and planet formation because only a small fraction of stars is found in high-density regions.  In the present study we will discuss the effect of different cluster density profiles, the dependency on the sample selection and the influence of observational constraints on the obtained results. We will demonstrate that local surface density measurements are rather limited in their ability to determine different star formation modes due to superposition effects.  Therefore the question whether dominant modes of clustered star formation exist in the solar neighbourhood is still open. ",
        "conclusions": "We investigated under which circumstances categorical distributions of local surface densities of young stellar objects -- here shortly referred to as \\emph{surface density plots} -- are suitable tools for investigating modes of clustered star formation and the dynamical influence of the star cluster environment on star and planet formation. Using different types of model star clusters we demonstrate how sensitive the results depend on the actual cluster density profile. Whereas for  narrow (for example Plummer-shaped) density distributions discrete cluster modes are easily identified as multiple peaks in the surface density plot; this is often not the case for distributions that span over a wider density range - for, example, concentrated King-type density distributions. Our findings imply that surface density plots of star-forming regions will not show multiple peaks unless the median density of the individual cluster modes differs by more than a factor of $\\sim$65.  The relative population size plays as well a role. Only if they do not differ by more than a factor of 5 the detection of discrete modes in the surface density plot is possible. Even, if one constructs equal sized samples, there might arise difficulties. If one combines different low-mass clusters to a single sample, it is very difficult to garantuee that they are in the same evolutionary stage. The cluster age is at least for embedded clusters not a reliable indicator for their dynamical stage.  The reason is that if star formation is ongoing and accelerated then averaging will always lead to approximately the same mean cluster age of $\\sim$1-2Myr. So a cluster just starting to form stars and one that has nearly finished the star formation process will both be attributed the same age. However, during that phase the cluster size, profile and surface density evolves considerably \\citep{2009A&A...498L..37P,2011A&A...536A..90P,2012Parmentier}. Including such different clusters in the same sample would lead to erroneous results. This means that the right sample choice is vital to determine whether a dominant mode of clustered star formation exists.  This means  that although one can conclude from multiple peaks in the surface density plot on the existance of discrete modes of clustered star formation, the reverse is not possible. We point out that unlike assumed in recent publications \\citep[e.g.][]{2010MNRAS.409L..54B} a smooth surface density plot does {\\em not} rule out the existence of dominant modes of clustered star formation. We thus caution against the use of surface density plots to determine whether dominant modes of clustered star formation exist.  However, surface density plots are potentially very useful in determining the dynamical influence of the cluster environment on star and planet formation. Yet a robust estimate requires high-resolution observations of rich star clusters to map the entire stellar population. Here we demonstrated that excluding regions of high local surface density in rich star clusters (like in Bressert et al. 2010) leads to underestimating the average local surface density {\\em not} as estimated in their study by at most a factor of two but by up to more than an order of magnitude.  Observations with instruments other than Spitzer (such as HST) are important for determining high surface density regions in  such clusters.  Another limitation that biases our understanding of star formation modes arises from restrictions of observational samples to the solar neighbourhood. Although there are good reasons for this approach such as sample completeness, one has to be aware that these results cannot be generalized to the Galaxy as for example, starburst clusters with their mostly much higher local surface densities are excluded.  Similarly, the age of the clusters included in the sample is an important factor. Dynamical interactions and stellar evolution in star clusters induce cluster expansion and hence act to lower their median local surface density with time.  This effect becomes even more pronounced if gas expulsion is taken into account \\citep[see e.g. the review of][]{2010RSPTA.368..829V}. Hence using a sample with a spread of cluster ages leads to an underestimate of the median local surface density of a given mode. This is of particular relevance for low-mass clusters that are expected to dissolve faster due to their short relaxation time.    In summary, a consistent analysis of the modes of clustered star formation requires a sample of isochronal clusters unlimited in mass and the development of a tool suitable to reveal potential discrete modes."
    },
    "1208/1208.2530_arXiv.txt": {
        "abstract": "We present a highly complete and reliable mid-infrared (MIR) colour selection of luminous AGN candidates using the 3.4, 4.6, and 12 \\mic bands of the \\wise survey.  The MIR colour wedge was defined using the wide-angle Bright Ultra-Hard \\xmm Survey ({\\tt BUXS}), one of the largest complete flux-limited samples of bright (${\\tt {\\it f}_{4.5-10\\,keV} > 6\\,x\\,10^{-14} \\flux}$) ``ultra-hard'' (4.5-10 keV) X-ray selected AGN to date.  {\\tt BUXS} includes 258 objects detected over a total sky area of 44.43 deg$^2$ of which 251 are spectroscopically identified and classified, with 145 being type-1 AGN and 106 type-2 AGN. Our technique is designed to select objects with red MIR power-law spectral energy distributions (SED) in the three shortest bands of \\wise and properly accounts for the errors in the photometry and deviations of the MIR SEDs from a pure power-law. The completeness of the MIR selection is a strong function of luminosity. At ${\\tt {\\it L}_{2-10\\,keV}>10^{44}\\lum}$, where the AGN is expected to dominate the MIR emission, $97.1_{-4.8}^{+2.2}\\%$ and $76.5_{-18.4}^{+13.3}\\%$ of the \\buxs type-1 and type-2 AGN meet the selection. Our technique shows one of the highest reliability and efficiency of detection of the X-ray selected luminous AGN population with \\wise amongst those in the literature. In the area covered by the \\buxs survey our selection identifies 2755 AGN candidates detected with \\sn$\\geq$5 in the three shorter wavelength bands of {\\rm WISE} with 38.5\\% having a detection at 2-10 keV X-ray energies. We also analyzed the possibility of including the 22{$\\mu$m} \\wise band to select AGN candidates, but neither the completeness nor the reliability of the selection improves. This is likely due to both the significantly shallower depth at 22\\mic compared with the first three bands of \\wise and star-formation contributing to the 22\\mic emission at the \\wise 22\\mic sensitivity. ",
        "introduction": "There is strong observational evidence that active galactic nuclei (AGN) play an important role in the formation and growth of galaxies (e.g. \\citealt{magorrian98}). Most supermassive black hole growth takes place during an obscured quasar phase, as suggested by the integrated energy density of the cosmic X-ray background \\citep{fabian99}. To understand the evolution of galaxies and to trace the energy output due to accretion and its cosmological evolution, it is critical to map the history of obscured accretion.  X-ray surveys with \\xmm and {\\textit{Chandra}} at energies $<$10 keV are sensitive to all but the most heavily obscured AGN (e.g. \\citealt{ceca08}). In Compton-thick AGN (rest-frame column densities exceeding ${\\rm N_H\\simeq 1.5\\times10^{24}\\,cm^{-2}}$) the observed flux below 10 keV can be as low as a few \\% of the intrinsic nuclear flux. In the Compton-thick regime the high energy photons that survive the photoelectric absorption get scattered in the absorber losing part of their energy (Compton down-scattering). This is an important effect that can significantly suppress the transmitted continuum (\\citealt{matt02}; \\citealt{murphy09}; \\citealt{yaqoob10}). The ongoing Swift/BAT and INTEGRAL/IBIS all-sky surveys at energies 15-200 keV are providing the least biased samples of absorbed AGN in the local Universe (e.g. \\citealt{bird07}; \\citealt{tueller08}; \\citealt{winter09}; \\citealt{burlon11}). However, even these surveys are biased against the most heavily absorbed Compton-thick AGN \\citep{burlon11}.  Surveys at mid-infrared (hereafter MIR) wavelengths ($\\gtrsim$5{$\\mu$m}) are much less affected by extinction since the obscuring dust re-emits the nuclear optical-to-X-ray radiation at infrared wavelengths. Clumpy torus models predict nearly isotropic emission in the MIR at wavelengths $\\gtrsim$12{$\\mu$m} (\\citealt{nenkova08}). Thus, MIR-based surveys (or the combination of MIR and data at shorter wavelengths) can potentially trace the elusive obscured accretion missed by hard X-ray surveys (e.g. \\citealt{daddi07}; \\citealt{fiore08}; \\citealt{georgantopoulos08}; \\citealt{fiore09}; \\citealt{severgnini12}). For example, it has been claimed that objects showing excess emission at $\\sim$24\\mic over that expected from star formation, termed \"infrared-excess galaxies\", might host heavily obscured and Compton-thick AGN (e.g. \\citealt{fiore08}; \\citealt{fiore09}). However the exact contribution of heavily obscured AGN to the infrared-excess galaxy population remains an open issue (e.g. \\citealt{alexander11}). Several MIR-based AGN selection techniques have been developed with data from the Spitzer Space Telescope Infrared Array Camera (IRAC; \\citealt{fazio04}) using colours and power-law selection (\\citealt{lacy04}; \\citealt{stern05}; \\citealt{alonso06}; \\citealt{donley08,donley12}). These techniques are very effective and reliable.  Galaxies dominated by AGN emission typically exhibit a characteristic red power-law spectral energy distribution (SED) in the MIR ($f_\\nu\\propto\\nu^{\\alpha}$ with $\\alpha\\leq$$-$0.5; \\citealt{alonso06}).  Thus, MIR power-law selection provides the cleanest samples of luminous AGN (e.g. \\citealt{donley08}). However, this technique is very sensitive to the reliability of the estimated photometric errors \\citep{donley12}.  The Wide-field Infrared Survey Explorer ({\\rm WISE}) has now completed the first sensitive ($\\sim$100-1000$\\times$ deeper than {\\rm IRAS}) coverage of the entire sky in the MIR\\footnote{https://ceres.ipac.caltech.edu/}\\citep{wright10}. Several colour-based regions, aimed at identifying luminous AGN, have already been proposed. These works have shown that \\wise can robustly separate AGN from normal galaxies and stars (e.g. \\citealt{assef10}; \\citealt{jarrett11}; \\citealt{stern12}). \\wise will be extremely efficient in identifying the rare highly luminous AGN up to the crucial epoch when the accretion power of the Universe peaked (\\red$\\sim$1-2). The all-sky \\wise survey will complement the deep Spitzer surveys, aimed to characterize the accretion phenomenon in the distant Universe.  This paper presents a highly reliable and complete MIR-based colour selection of AGN with {\\rm WISE}. Our technique is designed to select objects with red MIR power-law SEDs and properly accounts for the estimated typical errors in the photometry and deviations of the MIR SEDs from a pure power-law. The AGN wedge is defined using the wide-angle Bright Ultra-hard \\xmm Survey ({\\tt BUXS}; Mateos et al.\\ 2012c, in preparation). This survey is one of the largest complete flux-limited samples of bright ``ultra-hard'' (4.5-10 keV) X-ray selected AGN to date. Surveys such as \\buxs are extremely efficient in selecting AGN bright enough for reliable optical identifications and for detailed studies of their properties and evolution (e.g. HBS28, \\citealt{caccianiga04}; HBSS, \\citealt{ceca08}). \\buxs covers the region of the AGN redshift-luminosity parameter space that \\wise will sample. Thus, \\buxs offers a unique opportunity to define a highly complete and reliable MIR-based AGN selection with {\\rm WISE}. Thanks to the optical spectroscopic identifications available for $\\sim$97\\% of the \\buxs objects, and the high quality X-ray spectra, we have maximized the completeness of our MIR selection without compromising its reliability. In a forthcoming paper we will present and discuss the main properties of the optical/near-IR/MIR SEDs of the AGN in \\buxs (Mateos et al.\\ 2012c, in preparation).  This paper is organized as follows. Sections 2 and 3 briefly summarize the data sets. In Section 4 we present our MIR selection of AGN candidates using the three shorter wavelength bands of \\wise and the complete four bands, respectively and we discuss the completeness of the selection. We show the reliability of our AGN selection in Section 5. The results are summarized in Section 6. Throughout this paper errors are 90\\% confidence for a single parameter and we assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0={\\rm 70\\,km\\,s^{-1}\\,Mpc^{-1}}$. ",
        "conclusions": "We present a MIR power-law based selection of luminous AGN candidates using the 3.4, 4.6, and 12 \\mic bands of the \\wise survey. We defined an AGN wedge in the {\\tt log}(${\\tt {\\it f}_{4.6}/{\\it f}_{3.4}}$) vs. {\\tt log}(${\\tt {\\it f}_{12}/{\\it f}_{4.6}}$) colour-colour diagram using the Bright Ultra-Hard \\xmm Survey ({\\texttt{BUXS}}). This is one of the largest complete flux-limited samples of bright (${\\tt {\\it f}_{4.5-10\\,keV} > 6\\,x\\,10^{-14} \\flux}$) ``ultra-hard'' (4.5-10 keV) X-ray selected AGN to date. \\buxs includes 258 objects detected over a total sky area of 44.43 deg$^2$: 251 (97.3\\%) are spectroscopically identified and classified, with 145 being type-1 AGN and 106 type-2 AGN. Our technique is based on a MIR power-law selection and properly accounts for the errors in the photometry and deviations of the MIR spectral energy distributions from a pure power-law.  In flux-limited X-ray surveys, such as {\\texttt{BUXS}}, type-2 AGN are intrinsically less luminous than type-1 AGN. Thus, due to the strong dependence of the MIR selection completeness on the luminosity of the objects, a MIR AGN wedge necessarily picks out \\buxs type-1 AGN. However, at 2-10 keV luminosities above ${\\rm 10^{44}\\lum}$ the completeness of our MIR selection of type-1 and type-2 AGN is high and comparable for both types within the uncertainties. Our selection is highly complete at luminosities ${\\tt {\\it L}_{2-10\\,keV}>10^{44}\\lum}$ where our MIR wedge recovers $\\sim$97\\% and $\\sim$77\\% of the \\buxs type-1 and type-2 AGN, respectively. We identify 2755 AGN candidates in the 44.43 deg$^2$ \\buxs survey area of which 38.5\\% have detection in X-rays. In the \\buxs area where the X-ray observations have exposures $>$40\\,ks, the X-ray detection fraction rises to 49.8\\%. This is reasonable, as long X-ray exposures are required to detect intrinsically less luminous and/or heavily obscured AGN. A substantial fraction of the MIR AGN candidates remain undetected at 2-10\\,keV energies with the typical exposures in the {\\tt 2XMM} catalogue. These objects are the best candidates to account for the most heavily obscured/absorbed luminous AGN missed by hard X-ray surveys. Assuming that a 2-10 keV X-ray detection is a good tracer of AGN activity we demonstrate that our \\wise selection shows one of the highest reliability amongst those in the literature. This is crucial to obtain a clean MIR selection of powerful AGN. Furthermore, going down to a \\sn$\\geq$5 limit in the \\wise flux densities, we substantially increase the efficiency of detection of AGN with the reddest MIR colours. We also investigate a \\wise four-band AGN selection. We show, however, that by including the 22\\mic \\wise band neither the completeness nor the reliability of the selection improves. This is likely due to both the significantly shallower depth at 22\\mic compared with the first three bands of \\wise and star-formation contributing to the 22\\mic emission at the \\wise 22\\mic sensitivity. \\\\"
    },
    "1208/1208.0466.txt": {
        "abstract": "Over the past decades open clusters have been the subject of many studies. Such studies are crucial considering that the universality of the Initial Mass Function is still a subject of current investigations. Praesepe is an interesting open cluster for the study of the stellar and substellar mass function (MF), considering its intermediate age and its nearby distance. Here we present the results of a wide field, near--infrared study of Praesepe using the Data Release 9 (DR9) of the UKIRT Infrared Deep Sky Survey (UKIDSS) Galactic Clusters Survey (GCS). We obtained cluster candidates of Praesepe based on a 3$\\sigma$ astrometric and 5 band photometric selection. We derived a binary frequency for Praesepe of 25.6$\\pm$3.0\\% in the 0.2--0.45\\,M$_\\odot$ mass range, 19.6$\\pm$3.0\\% for 0.1--0.2\\,M$_\\odot$, and 23.2$\\pm$5.6\\% for 0.07--0.1\\,M$_\\odot$. We also studied the variability of the cluster candidates of Praesepe and we conclude that seven objects could be variable. We inferred the luminosity function of Praesepe in the $Z-$ and $J-$ bands and derived its MF. We observe that our determination of the MF of Praesepe differs from previous studies: while previous MFs present an increase from 0.6 to 0.1\\,M$_\\odot$, our MF shows a decrease. We looked at the MF of Praesepe in two different regions of the cluster, i.e.~within and beyond 1.25$^{\\circ}$, and we observed that both regions present a MF which decrease to lower masses. We compared our results with the Hyades, the Pleiades and $\\alpha$~Per MF in the mass range of 0.072--0.6\\,M$_\\odot$ and showed that the Praesepe MF is more similar to $\\alpha$~Per although they are respectively aged $\\sim$85 and $\\sim$600\\,Myr. Even though of similar age, the Praesepe remains different than the Hyades, with a decrease in the MF of only $\\sim$0.2\\,dex from 0.6 down to 0.1\\,$M_{\\odot}$, compared to $\\sim$1\\,dex for the Hyades. ",
        "introduction": "\\label{intro}  Over the past decades, open clusters have been the subject of many studies \\citep[e.g.][and references therein]{bastian2010}.  Such studies have brought new insights into brown dwarf formation \\citep[e.g.][]{kumar2007,boudreault2009,bejar2011}, on the discovery of young L and T dwarf and free-floating planets \\citep[e.g.][]{bouvier2008,lodieu2008,zo2008,bihain2009,lucas2010,quanz2010,pr2011}, and on our understanding of the stellar/substellar mass function (MF) \\citep[see review by][]{bastian2010} and their populations in the Galactic field and in open clusters \\citep[see review by][]{chabrier2003}.  Such studies are crucial considering that the universality of the Initial Mass Function (IMF) is still a subject of current investigations \\citep[e.g.][]{kroupa2002,covey2011,myers2011,leigh2012,marks2012}. Most works on the substellar MF have focused on young open clusters with ages less than $\\sim$150\\,Myr. This is partly because brown dwarfs (BDs) are bright when they are young, thus aiding the detection of the least massive objects.  However, the extension of MF studies to older clusters is vital as it allows us to study the intrinsic evolution of BDs and how the stellar and substellar population itself evolves.  Praesepe (M\\,44, NGC\\,2632, RA=8$^{h}$40.4$^{m}$, DEC=$+$19$^{\\circ}$41$'$) is an interesting open cluster to study the MF in the stellar and substellar regimes, considering its age \\citep[$\\tau=590^{+150}_{-120}$\\,Myr;][]{fossati2008} and its distance \\citep[$(m-M)_{0}=6.30\\pm0.07$\\,mag, $d=181.97^{+5.96}_{-5.77}$\\,pc;][]{leeuwen2009}, its known proper motion \\citep[$\\mu_{\\alpha}=-35.81\\pm0.29$\\,mas/yr and $\\mu_{\\delta}=-12.85\\pm0.24$\\,mas/yr;][]{leeuwen2009}, and the low extinction towards this cluster \\citep[$E(B-V)=0.027\\pm0.004$;][]{taylor2006}, while determinations of the metallicity of Praesepe go from solar-type with [Fe/H]$=0.038\\pm$0.039\\,dex \\citep{friel1992} to slightly metal rich with [Fe/H]$=0.27\\pm$0.10\\,dex \\citep{pace2008}. So far, several surveys for stellar and substellar objects in the open cluster Praesepe have been performed \\citep[e.g.][]{jones1983,hodgkin1999,baker2010}. Some surveys have the advantage of covering a large area and using proper motions, but are rather shallow \\citep[e.g.][]{hambly1995,kraus2007}, while other surveys used deep photometry, but lacked wide areal coverage \\citep[e.g.][]{gg2006,boudreault2010,wang2011}.  The UKIRT Infrared Deep Sky Survey \\citep[UKIDSS;][]{lawrence2007} is a deep large scale infrared survey conducted with the wide field camera WFCAM \\citep{casali07} on UKIRT (Mauna Kea, Hawai'i). The survey is subdivided into five components: the Large Area Survey, the Galactic Clusters Survey (hereafter GCS), the Galactic Plane Survey, the Deep Extragalactic Survey, and the Ultra-Deep Survey. The GCS aims at covering $\\sim$1000 square degrees in 10 star forming regions and open clusters down to $K$ = 18.4 mag at two epochs. The main scientific driver of the survey is to study the IMF and its dependence with environment in the substellar regime using an homogeneous set of low-mass stars and brown dwarfs over large areas in several regions. The UKIDSS GCS is, therefore, a perfect tool to study the open cluster Praesepe, considering the large coverage from the UKIDSS Data Release 9 (DR9) with its relative deep photometry spanning from $J$$\\sim$10.9\\,mag (i.e.~$\\sim$0.7\\,M$_\\odot$) down to $J$$\\sim$19.3\\,mag (i.e.~$\\sim$55\\,M$_{\\rm Jup}$), combined with astrometric information.  Here we present the results of a wide--field near--infrared study of the Praesepe cluster using the DR9 of the UKIDSS GCS. The paper is structured as follows. First we present the dataset used in our analysis (Section \\ref{phot-and-astrom}), followed by a cross--match with previous surveys (Section \\ref{cross-match}). Then we extract the new stellar and substellar members in Praesepe based on our selection criteria (Section \\ref{new-substellar-members}). We discuss the level of contamination (Section \\ref{contamination}), the multiplicity of low-mass Praesepe members (Section \\ref{binarity}), and the variability of our cluster candidates (Section \\ref{variability}). Finally we derive the luminosity function (LF) and MF of Praesepe (Section \\ref{lf-and-mf}). ",
        "conclusions": "\\label{conclusion}  In this paper we presented the results of a wide field, near--infrared study of the Praesepe cluster using the DR9 of the UKIRT Infrared Deep Sky Survey Galactic Clusters Survey. We performed an astrometric and photometric selection of 1,116 cluster candidates out of the 218,141 point sources detected towards Praesepe.  Possible sources of contamination include Galactic disk late-type and giant stars and unresolved galaxies.  We estimate a contamination of 11.9\\,\\% above 0.4\\,M$_\\odot$, 9.8\\,\\% in the mass range 0.15--0.4\\,M$_\\odot$, and 23.8\\,\\% below 0.15\\,M$_\\odot$.  We investigated the binary frequency of Praesepe using the photometry and colours from our cluster candidates. We observe a binary fraction similar to the simulation of \\citet{bate2012} between 0.07--0.1\\,M$_\\odot$, $\\sim$1.5$\\sigma$ difference in the 0.2--0.45\\,M$_\\odot$ mass interval, and significantly lower by more than 3$\\sigma$ for the mass range 0.1--0.2\\,M$_\\odot$. On the other hand, the binary fraction from \\citet{pinfield2003} are higher than our values and those of \\citet{bate2012}. However, we note that two other works focusing on field low-mass stars have also derived binary fractions lower than \\citet{bate2012}.  We also studied the variability of the Praesepe candidates using the two $K$--band epochs provided by the GCS. We identified seven candidate variables, including one in the substellar regime.  We derived the luminosity function of Praesepe in $Z$ and $J$--band here. We observed that the peak of the $J$--band luminosity function is one magnitude brighter than the one reported by \\citet{boudreault2010}.  Finally, we determined the mass function of Praesepe, which differs from previous studies: while previous MFs showed an increase from 0.6 to 0.1\\,M$_\\odot$, our MF shows a decrease. We looked at the MF of Praesepe at two different regions of the cluster, i.e.~within and beyond 1.25$^{\\circ}$, and we observed that both regions show an MF which decreases to lower masses. We compared our MF of Praesepe in the mass range 0.072--0.6\\,M$_\\odot$ with the ones of the Hyades, the Pleiades and $\\alpha$~Per. We concluded that our MF of Praesepe is most similar to the MF of $\\alpha$~Per although they are respectively of $\\sim$85 and $\\sim$600\\,Myr. Even though of similar age, the Praesepe appears different to the Hyades, with a decrease in the MF of only $\\sim$0.2\\,dex from 0.6 down to 0.1\\,$M_{\\odot}$, compared to $\\sim$1\\,dex for the Hyades."
    },
    "1208/1208.0229_arXiv.txt": {
        "abstract": "{ It is often assumed that the halo-patch fluctuation field can be written as a Taylor series in the initial Lagrangian dark matter density fluctuation field.  We show that if this Lagrangian bias is local, and the initial conditions are Gaussian, then the two-point cross-correlation between halos and mass should be linearly proportional to the mass-mass auto-correlation function. This statement is exact and valid on all scales; there are no higher order contributions, e.g., from terms proportional to products or convolutions of two-point functions, which one might have thought would appear upon truncating the Taylor series of the halo bias function.  In addition, the auto-correlation function of locally biased tracers can be written as a Taylor series in the auto-correlation function of the mass; there are no terms involving, e.g., derivatives or convolutions. Moreover, although the leading order coefficient, the linear bias factor of the auto-correlation function is just the square of that for the cross-correlation, it is the same as that obtained from expanding the mean number of halos as a function of the local density only in the large-scale limit.  In principle, these relations allow simple tests of whether or not halo bias is indeed local in Lagrangian space.  We discuss why things are more complicated in practice.  We also discuss our results in light of recent work on the renormalizability of halo bias, demonstrating that it is better to renormalize than not.  We use the Lognormal model to illustrate many of our findings.} ",
        "introduction": "Galaxies are biased tracers of the dark matter distribution \\cite{nk1984, bbks1986}.  In the local bias model, the abundance of the biased tracers at a given position is assumed to be related to the mass at the same position.  The simplest version of this model, in which the smoothed overdensity field of the biased tracers $\\delta_b$ is treated as though it were a deterministic function of the (similarly smoothed) real-space mass overdensity field $\\delta_m$, has been the subject of much study \\cite{as1988, pc1993, sw1998}.  Following ref.~\\cite{fg1993}, it has become common to write the local model as \\begin{equation} \\delta_b = f(\\delta_m) = \\sum_{i>0} \\frac{b_i}{i!} \\,(\\delta_m^i - \\langle\\delta_m^i\\rangle), \\label{lb} \\end{equation} where $b_i$ is the bias coefficient of order $i$.  Note that this model ensures $\\langle\\delta_b\\rangle = 0$ by subtracting-off the $\\langle\\delta_m^i\\rangle$ terms.\\\\ One of the points we make in what follows is that one should, instead, normalize using a multiplicative factor which ensures that $\\langle 1+\\delta_b\\rangle = 1$.  I.e., given an expansion of the form~(\\ref{lb}), one should always work with \\begin{equation} \\delta_{B} \\equiv \\frac{1+\\delta_{b} - \\langle 1+\\delta_b\\rangle} {\\langle 1+\\delta_b\\rangle}. \\label{lbcorrect} \\end{equation}  Although the model was invoked to describe the bias with respect to the late-time nonlinear Eulerian field $\\delta_m^E$, since it is only invoked on large scales, it is often assumed, and sometimes explicitly used, to describe the bias with respect to the initial Lagrangian field $\\delta_m^L$ (on large scales, these are expected to be similar).   However, the two best studied models of bias with respect to the Lagrangian field -- those associated with peaks, and patches which form halos, in a Gaussian random field -- behave rather differently from naive expectations based on the local bias model.  In particular, in the local bias model calculation of the cross-correlation between the biased tracers and the initial field, $\\langle\\delta_b\\delta_m\\rangle$, one proceeds by writing the bias function as a Taylor series, and then expanding order by order in $\\delta_m$.  This means that one expects higher-order terms to contribute.  However, for peaks and in Gaussian initial conditions, the exact answer is $\\langle\\delta_b\\delta_m\\rangle = b_1\\,\\langle\\delta_m^2\\rangle$ \\cite{dcss2010}.  Therefore, in the expansion referred to above, all the higher-order terms must cancel out.  We show below that this is also true for another popular choice of the bias -- a Lognormal mapping between the biased field and $\\delta_m$ -- and then that this is generically true in models where the bias is local and deterministic with respect to the initial Gaussian random field.  I.e., the cross-correlation is always only linearly proportional to the auto-correlation signal of the dark matter, although, in general, the linear bias factor need not equal $b_1$ of the Taylor series.  We then show that the auto-correlation function of the biased tracers can always be written as a Taylor series in the auto-correlation function of the original (unbiased) mass fluctuation field; derivatives and convolutions do not enter.  Along the way, our analysis connects to recent work on renormalized bias \\cite{pm2006}, showing, e.g., that this renormalization is required if the Lognormal mapping is to make sensible predictions.  We have structured the discussion as follows. In section~\\ref{ln} we discuss the Lognormal mapping, since it turns out that all quantities of interest can be computed exactly (no truncation of the sums is required).  These exact expressions exhibit some curious properties which are not obvious if one trucates the sums.  Then we show that how one normalizes such purely formal expansions plays an important role, and why eq.~(\\ref{lbcorrect}) is to be preferred over eq.~(\\ref{lb}). Section~\\ref{pks} explores a more complex example in which some of this simplicity is lost, before showing the general result in section~\\ref{general}. Section~\\ref{renorm} shows how to proceed if the full expansion is not available, so one is constrained to work with a truncated series.  Section~\\ref{hermiteBias} connects this analysis to some of the earliest work on this subject (ref.~\\cite{as1988}), discussing how the coefficients of the Taylor series expansion of $\\delta_B$ in terms of $\\delta_m$ are related to those obtained from expanding $\\xi_{BB}$ in terms of $\\xi_{mm}$.  Section~\\ref{halos} discusses halo bias in the context of these results, and section~\\ref{porq} revisits a technical point about halo bias, first made by ref.~\\cite{sl1999} but often overlooked, which complicates the use of cross-correlations for testing the hypothesis that Lagrangian halo bias is indeed local, a subject of much recent interest \\cite{mg2011, psp2012, css2012, mps2012}.  A final section summarizes. ",
        "conclusions": "In local bias models it is assumed that the biased field $1+\\delta_b$ can be written as a function of the underlying density field at the same position.  We showed that, if the underlying field is Gaussian, then the cross-correlation between the biased field and the original one is linearly proportional to the auto-correlation function of the original field.  This is an exact result, valid on all scales, and is not a consequence of truncating expansions, etc., as most previous treatments assume.  While this is implicit in some previous work (e.g.~ref.~\\cite{mss2010}), it has not been highlighted before.  If one has been careful to ensure that $\\langle 1+\\delta_b\\rangle = 1$ (by using eq.~(\\ref{lbcorrect}) rather than eq.~(\\ref{lb})), then the constant of proportionality is easily related to the first coefficient of the Taylor series expansion of the biased field, although they are not equal in general (eq.~(\\ref{Xcorrect})).  In addition, we showed that, to leading order, the ratio of the square of the cross-correlation to the auto-correlation of the bias tracers equals the correlation function of the underlying field (eqs.~(\\ref{Acorrect}) and~(\\ref{ratio2xi})).  We also explored the consequences of truncating these expansions, demonstrating that it is better to renormalize all truncated expansions (eq.~(\\ref{deltabj})) than to not.  In this respect, our results agree with ref.~\\cite{pm2006}.  Indeed, although our work has concentrated on Lagrangian bias, the multiplicative normalization (our eq.~(\\ref{lbcorrect})) also holds also for Eulerian mass field, although providing an explicit expression when the mass field is not gaussian is more complicated (see ref.~\\cite{pm2006} for the implications in Fourier space).  Expanding the biased field using Hermite polynomials rather than powers of the overdensity, as was done by ref.~\\cite{as1988}, provided a easy way to see a number of our results for local bias in the initial, Lagrangian, Gaussian field.  Although the coefficients $B_k$ of this expansion are functions of the scale on which the bias transformation is applied (eq.~(\\ref{Jkszalay})), for the Lognormal distribution, as well as for excursion set halos, the scale dependence of these coefficients is trivial:  $B_k = b_k\\sigma_{\\rm L}^k$ (eq.~(\\ref{Jkhalos})), where $b_k$ are the scale-independent peak-background split bias factors.  (The Lognormal has $b_k=b^k$, whereas excursion set halos have $b_k$ given by eq.~(\\ref{bkhalos}).)  Therefore, if one expands the (Lagrangian space) halo auto-correlation function in powers of the mass correlation, then the coefficient of the $k$th order term in the expansion is simply $b_k^2$ (eq.~(\\ref{szalayXi})).  These coefficients for the expansion of $1 + \\xi_{hh}$ in powers of $\\xi_{LL}$ are independent of scale $L$, even though the bias coefficients in the Taylor expansion of the field $1+\\delta_h(\\delta_L)$ itself do `run' with $L$.  This property of halo bias in Lagrangian space has not been emphasized before.  The `running' of the $1+\\delta_h(\\delta_L)$ bias factors is most easily understood by noting that the Hermite polynomials arise naturally if one phrases the question of local bias in Fourier rather than real space \\cite{tm2011}, a connection made recently in ref.~\\cite[in particular, their Appendix~B]{mps2012}.  Moreover, ref.~\\cite{tm2011} notes that it is better to work with `renormalized' parameters $c_n$ rather than the original bias parameters $b_n$ (but see \\cite{mps2012} for how this renormalization should actually be done).  In effect, this corresponds to working with our eq.~(\\ref{lbcorrect}) rather than eq.~(\\ref{lb}); our analysis shows why this is necessary.  This connection to Fourier space bias is rich, because we showed that the simple relation $B_k = b_k\\sigma_{\\rm L}^k$ is not generic.  For the peaks transformation of eq.~(\\ref{bpkapprox}), both $b_k$ and $B_k$ `run' with $L$.  Nevertheless, the $B_k$ satisfy the same relations between the rescaled values that the original $b_k$ (in the peak background split limit) do.  Exploring whether this is generic is the subject of ongoing work.  Unfortunately, although the analysis based on Hermite polynomials is formally correct (and elegant), there are two reasons why, at least for describing halos, it cannot be valid for all smoothing scales $L$ and separations $r$.  First, halos do not overlap; this makes $\\xi_{hh}$ of the halo point process tend to $-1$ on scales smaller than $\\sim R_h$ \\cite{mw1996}.  As a result, this halo exclusion limits the range in $r$ over which the local model can be applied.  The second is that, at least in the excursion set definition of halos, there is a nonlocal requirement on the density field:  this was the subject of section~\\ref{porq}, which argued that this modifies the pdf over which one should compute ensemble averages.  Accounting for this makes the ratio of the halo-mass cross-correlation bias scale dependent (our eq.~(\\ref{b1q}), which is eq. (17) of \\cite{sl1999}).  The small scale limiting value of this modified expression yields $\\delta_c$, which is the correct `one-halo' contribution to the Lagrangian space cross-correlation (the Lagrangian overdensity within a region which is destined to form a halo equals $\\delta_c$ by definition), a fact which has not been emphasized before.  This suggests that averaging over the more appropriate pdf not only leads to sensible results, but accounts for halo exclusion as well, so it may be worth exploring further.  This may be particularly interesting because, in the large separation limit, it leads to $\\xi_{hh} - b_\\times^2\\,\\xi_{mm} < 0$ (eq.~(27) of \\cite{sl1999}).  There is a third reason why, at least for describing halos, the local bias model is unlikely to be valid for all smoothing scales $L$ and separations $r$.  This is related to the fact that the tidal field influences halo formation \\cite{smt2001}.  The correlation with the tidal field leads, generically, to nonlocal bias even in Lagrangian space.  This is explored further in ref.~\\cite{scs2012}, where the nonlocal bias terms are shown to matter most for massive halos, though the nonlocal effects are subdominant on large scales.  If the gravitationally evolved nonlinear Eulerian mass field was a locally biased version of the initial field (it is not), then expanding in Hermite polynomials would also be the prefered way of describing halo bias (assuming the initial conditions were Gaussian).  This is because local Lagrangian bias, with local nonlinear evolution, leads to local bias with respect to the Eulerian field.  Conversely local bias with respect to the Eulerian field could be mapped back to local Lagrangian bias (with different bias factors, determined by the local linear-nonlinear mapping).  Some of the scalings which are characteristic of local Lagrangian bias will survive in Eulerian bias as well.  For example, the Lognormal transformation has one free parameter $b$.  If we use $b'$ to model the nonlinear mass field, and another $b$ to model the biased tracers (incidentally, this means that one should think of the biased field $1+\\delta_b$ as being the nonlinear field $1+\\delta_{b'}$ raised to the $b/b'$ power), then we may interpret our eq.~(\\ref{xinm}) as the Eulerian space cross-correlation between the biased field and the nonlinear mass field.  Since $\\langle\\delta_b\\delta_L|r\\rangle$ and $\\langle\\delta_{b'}\\delta_L|r\\rangle$ equal $b\\,\\xi_{LL}(r)$ and $b'\\,\\xi_{LL}(r)$ for all $r$, one might naively have thought that $\\langle\\delta_b\\delta_{b'}|r\\rangle$ would also be linearly proportional to $\\xi_{LL}(r)$.  Not only is this not true, eq.~(\\ref{xinm}) shows that it is not proportional to $\\xi_{b'b'}(r)$ either. In particular, eq.~(\\ref{xinm}) shows that in this model of local Eulerian bias, the cross-correlation is linearly proportional to the auto-correlation of the mass only when $\\xi_{LL}\\ll 1$ (i.e., for sufficiently large $r$).   See Ref.~\\cite{sw1998} for more discussion of this limit of local Eulerian bias.  Recent work has emphasized the fact that because nonlinear evolution is nonlocal, local Lagrangian bias will lead to nonlocal Eulerian bias, and vice versa \\cite{tm2011,psp2012,css2012,bsdm2012}.  The Hermite polynomials are the orthogonal polynomials associated with a Gaussian field.  Therefore, the usefulness of Szalay's work for local Lagrangian bias suggests that, if bias is local with respect to the nonlinear, non-Gaussian field, then it would be natural to write Eulerian bias using the orthogonal polynomials of this non-Gaussian field.  This is the subject of work in progress.  Finally, we note that our results depend only on the assumption that the Lagrangian matter field is Gaussian, so one might have thought that they also hold for modified gravity models.  However, for such models, $k$-dependence of the linear growth factor is generic, with the consequence that the assumption that Lagrangian halo bias is local is no longer so attractive \\cite{phs2011}."
    },
    "1208/1208.5583_arXiv.txt": {
        "abstract": "We present a catalog of high-velocity clouds in the region of the Magellanic Leading Arm. The catalog is based on neutral hydrogen (\\HI) observations from the Parkes Galactic All-Sky Survey (GASS). Excellent spectral resolution allows clouds with narrow-line components to be resolved. The total number of detected clouds is 419. We describe the method of cataloging and present the basic parameters of the clouds. We discuss the general distribution of the high-velocity clouds and classify the clouds based on their morphological type. The presence of a significant number of head-tail clouds and their distribution in the region is discussed in the context of Magellanic System simulations. We suggest that ram-pressure stripping is a more important factor than tidal forces for the morphology and formation of the Magellanic Leading Arm and that different environmental conditions might explain the morphological difference between the Magellanic Leading Arm and Magellanic Stream. We also discuss a newly identified population of clouds that forms the LA IV and a new diffuse bridge-like feature connecting the LA II and III complexes. ",
        "introduction": "}  Some atomic neutral hydrogen (\\HI) concentrations surrounding our Galaxy have anomalous velocities that are forbidden by a simple Galactic rotation model. These so-called anomalous-velocity clouds contain \\HI\\ without a stellar counterpart. They can be classified into two velocity based groups: intermediate velocity clouds (IVCs; \\citealp{MZ61,BT66}); high-velocity clouds (HVCs; \\citealp{Muller63}). The classification of IVCs and HVCs is based on the deviation velocity, which is defined as the smallest difference between the velocity of the cloud and the Galactic rotational velocity  \\citep{Wakker91}. High-velocity clouds are particularly interesting because they are thought to represent the flow of baryons in or out of the Galactic disk, which influences the formation and evolution of our Galaxy. Despite being important in the context of galaxy formation and evolution, their origin and physical characteristics are still under debate.  Possible explanations of the origin of HVCs can be traced back to an early study by \\citet{Oort66}. One of his hypotheses suggested that the HVCs have an extragalactic origin. This hypothesis received more recent support from \\citet{Blitz99} with the argument that HVCs are dark matter dominated clouds in the Local Group with distances of hundreds of kiloparsecs. A similar study by \\citet{BB99} also claimed that compact and isolated HVCs lie at extragalactic distances. Another popular HVC origin hypothesis is the Galactic fountain model, in which the gas is blown out of the disk by supernovae, cools and then rains back down (see \\eg, \\citealp{Houck90}). In this scenario, the fountain gas can rise as high as $\\sim$10~kpc above the disk \\citep{DA00}.  The previous large-scale surveys, the Leiden--Argentine--Bonn Galactic \\HI\\ survey (LAB; \\citealp{Kalberla05}) and the \\HI\\ Parkes All-Sky Survey (HIPASS; \\citealp{Barnes01}) have provided opportunities to study HVCs on a global scale to assist in the understanding of their origin and physical properties (see \\eg, \\citealp{WW91,Putman02}). LAB covered the entire sky with an angular resolution of 36$\\arcmin$ and a spectral resolution of 1.3~\\kms. HIPASS was conducted with a better angular resolution (16$\\arcmin$) but lower spectral resolution (18~\\kms). A comprehensive catalog of Southern HVCs based on the HIPASS data was presented in \\citet{Putman02}, hereafter P02. The catalog covers the high-velocity \\HI\\ sky south of declination +2\\degr\\ and within the Local Standard of Rest velocity ($V_{LSR}$) range of +500 to $-$500~\\kms. It provides the spatial and kinematic distributions as well as the properties of high-velocity clouds. Even though the P02 catalog includes a complete census of HVCs south of declination +2$\\degr$, the nature of the in-scan bandpass calibration technique filtered out some of large-scale structure of the Milky Way and the Magellanic System.  Among HVC complexes, the Magellanic System is the most interesting given that it is the only closest extragalactic gaseous stream to our Galaxy. The Magellanic System consists of a coherent gas stream originating from the Magellanic Clouds (MCs), \\ie, Magellanic Stream (MS) and Leading Arm (LA) \\citep{Mathewson74}. The MS is trailing the MCs and has a complex filamentary structure. On the other hand, the LA is clumpy and dominated by three distinctive large complexes, namely the LA I, LA II and LA III \\citep{Putman98,Bruns05}. An extended feature of the MS has recently been discovered by \\citet{Nidever10}, which reveals the total length of the MS as $\\sim$200\\degr\\ across the sky. Another recent report of several filaments that are aligned with the MS also suggests that MS is wider than previously thought \\citep{Westmeier11}.  The formation of the MS and LA is generally believed to have been caused by the tidal interaction between the Milky Way and Magellanic Clouds. Theoretical models with tidal stripping, gravitational and hydrodynamical interactions can reproduce global observed \\HI\\ column density and velocity distributions (see \\eg, \\citealp{Connors06,Mastropietro05}). However, these models do not provide a satisfactory explanation for the formation mechanism of the MS and LA. With the recent Hubble Space Telescope proper motion measurements of the MCs \\citep{K06b,K06a}, a new unbound orbit for the MCs with a first passage scenario was proposed by \\citet{Besla07}. The result is surprising given that the new orbit does not provide sufficient time for tidal and ram-pressure stripping mechanisms to produce the MS \\citep{Stanimirovic08}. To circumvent the problem raised by the first passage scenario, \\citet{Nidever08} proposed a new blowout hypothesis. They suggested that the supergiant shells in the dense southeast \\HI\\ overdensity region are blown out from the Large Magellanic Cloud (LMC) to larger radii where ram-pressure and/or tidal forces can be more easier to strip the gas and form the MS and LA. Nevertheless, with recent higher measurements of the Milky way's circular velocity (\\eg, 251~\\kms; \\citealp{Reid04}) compared to the IAU standard of 220~\\kms, there remains the distinct possibility that a multi-orbit scenario is plausible \\citep{SL09}. Recent multi-orbit simulations have also reproduced the observed structure of the MS, including bifurcation of the two filaments, whilst remaining consistent with the proper motion data \\citep{DB11,DB12}. However, none of the theoretical models to date have been able to accurately reproduce the observed structure of the LA.  To study the Magellanic System in detail, \\citet{Bruns05} carried out a narrow-band Parkes \\HI\\ survey. In contrast to HIPASS, which was not designed to accurately measure low-velocity Galactic \\HI\\ gas, the survey was designed exclusively to study the Magellanic System. The Br\\\"{u}ns' survey has a similar angular resolution and spectral resolution to the Galactic All-Sky Survey (GASS; \\citealp{MG09}; \\citealp{Kalberla10}) (see \\S2), but with limited sky coverage. Another high-resolution \\HI\\ study that concentrated on the northern tip of the MS was carried out by \\citet{Stanimirovic08}. This study was part of the Galactic studies with Arecibo $L$-band Feed Array (GALFA).  The current work utilizes the GASS data for studying the general distribution and morphological types of HVC in the region of the Magellanic Leading Arm. The GASS data have better sky coverage than the Br\\\"{u}ns' survey and higher spectral resolution than HIPASS. The study of HVCs in the vicinity of the LA gives us clues to understand: (1) the formation of the LA; (2) the physical properties of the HVCs; and most importantly, (3) the role of infalling gas in the context of galaxy evolution and formation. In \\S2, we describe the GASS data and the procedures for cloud search algorithms. We present the catalog and the general distribution of clouds in \\S3 and \\S4. Classification of the clouds and interpretation of the distribution for each group are given in \\S5. Finally, we report on the new extended features of the Magellanic Leading Arm and discuss the implications of HVCs for the formation and origin of the LA in \\S6. Conclusions are drawn in \\S7. ",
        "conclusions": "}  We have produced a catalog of high velocity clouds in the region of the Magellanic Leading Arm from Parkes Galactic All-Sky Survey data, using the cloud search algorithm {\\it Duchamp}.  We used {\\it Duchamp} to parametrize cloud properties including position, velocity and velocity FWHM.  We determined the angular size of sources via 2-dimensional Gaussian fitting and peak \\HI\\ column density via searching the brightest pixel in the integrated maps. Comparison between our HVC catalog with that of \\citet{Putman02} in the same region and velocity range shows that the high spectral resolution of GASS allows us to recover clouds with narrow line widths. The total number of detected HVCs is 448 for P02 catalog and 419 for our catalog. The combined catalog contains $\\sim$625 unique clouds.  We have presented the general distribution of HVCs in the catalog. The kinematic distributions with respect to Galactic longitude and latitude are generally consistent with the findings in P02. A trend of decreasing number of clouds from higher to lower Galactic latitude as velocity increases in all velocity reference frames was found. A morphological classification of clouds was presented, and distributions of each type were discussed.  An extended feature in the LA I complex that was not covered in the detailed study of the Magellanic System by \\citet{Bruns05} was noted. A new population of clouds that forms the LA IV and an extended feature that forms a diffuse ``bridge'' connecting the LA II and III complexes were also discovered. The discovery of the LA III extended feature demonstrates the importance of brightness temperature sensitivity and spectral resolution for an all-sky survey. Simulations have yet to reproduce this feature of LA.  The most significant result in this study was the detection of a large number of head-tail clouds in the region of the LA as compared to the MS, suggesting that ram-pressure stripping is relatively more important than gravitational forces for the morphology and formation of the LA. The LA I and II themselves are large head-tail clouds, which are moving toward higher Galactic latitudes and both show a large velocity gradient, with the head being lower than the tail. We found that there was no preferred pointing direction for the small head-tail clouds. This suggests a scenario where the clouds are produced in a turbulent flow where incoming warm neutral gas collides with the hot halo ISM. The cloud morphologies are strongly correlated to the degree of turbulence in the ISM. The presence of strong turbulence is probably the cause for the observed morphologies and properties of clouds in the region.  A dichotomy in velocity for the head-tail and symmetric HVCs above and below the Galactic plane was found. Since such dichotomy is also seen in simulations of \\citet{DB12}, we suggest that an orbital effect is the cause. Finally, using the typical \\HI\\ column density of cloud and tangential velocities above and below the Galactic plane, we infer a small difference in halo density. This suggests that the LA II and LA II+LA III are interacting with different halo environments, which might explain the morphological difference between them.   \\appendix"
    },
    "1208/1208.5898_arXiv.txt": {
        "abstract": "Results are presented from a 500~ks long XMM-Newton observation of the Narrow-Line Seyfert 1 galaxy IRAS\\,13224-3809. The source is rapidly variable on timescales down to a few 100~s. The spectrum shows strong broad Fe-K and L emission features which are interpreted as arising from reflection from the inner parts of an accretion disc around a rapidly spinning black hole. Assuming a power-law emissivity for the reflected flux and that the innermost radius corresponds to the innermost stable circular orbit, the black hole spin is measured to be 0.988 with a statistical precision better than one per cent. Systematic uncertainties are discussed. A soft X-ray lag of 100\\,s confirms this scenario. The bulk of the power-law continuum source is located at a radius of 2--3 gravitational radii. ",
        "introduction": "The X-ray emission from most Narrow-Line Seyfert 1 galaxies (NLS1) is characterised by a steep soft X-ray spectrum and rapid variability. The most extreme such objects are 1H\\,0707-495 and IRAS13228-3809, which both show a sharp drop above 7~keV in XMM Newton spectra (Boller et al 2002, 2003). 1H\\,0707-495 has been further studied several times with XMM, including a long 500~ks dataset in 2008 which revealed broad iron K and L lines and a soft lag of about 30~s (Fabian et al 2009; Zoghbi et al 2010, 2011). In contrast, IRAS13224-3809 only had 64~ks of XMM data (Ponti et al 2010; Gallo et al 2004), despite showing spectacular variability during ROSAT (Boller et al 1997) and ASCA (Dewangan et al 2002) observations. New observations totalling 500~ks have now been made with XMM in 2011 and reported here.  The unusual spectrum and 7~keV drop of both objects have been interpreted as due to either intervening absorption or strong relativistic blurring of a reflection component (Boller et al 2002, 2003; Fabian et al 2004).  1H\\,0707-495 dropped into a low state for about 2 months at the start of 2011 during which an XMM spectrum showed evidence for even more blurring. The results are consistent with the power-law component of the X-ray source lying within one gravitational radius of the central black hole (Fabian et al 2012). In the normal state, one third of this component extends to $\\sim 20r_{\\rm g}$.  The combination of the above results with the reverberation lags in 1H\\,0707-495 and in over a dozen other sources (Emanoulopolous, McHardy \\& Papdakis 2010; Tripathi et al 2011; De Marco et al 2011, 2012; Zoghbi \\& Fabian 2011 and Zoghbi et al 2012) provides very strong support for the reflection model for the X-ray emission of Seyfert galaxies. In this model the primary power-law component lies above the inner accretion disc around the black hole and produces the X-ray reflection component by irradiation of the disc (see e.g. Fabian \\& Ross 2010). The soft lags are then the light travel time difference between the power-law and reflection components as detected by the observer. The new data presented in this paper are interpreted within the reflection model.  IRAS13224-3809 is a radio quiet (1.4~GHz flux of 5.4~mJy, Feain et al 2009) NLS1 at redshift $z=0.066$. For a flat $\\Lambda$CDM cosmology with $H_0=71\\kmpspmp$, its luminosity distance is 293~Mpc. ",
        "conclusions": "IRAS\\,13224-3809 is remarkably similar in overall X-ray behaviour to 1H\\,0707-495. The variability of IRAS\\,13224-3809 may be the most extreme. We shall explore the behaviour of the source as a function of time and flux in more detail in later work.  The X-ray spectra of both sources require high iron abundance ($A_{\\rm Fe}\\sim 10-20$).  In recent work, Wang et al. (2012) have presented a strong correlation between metallicity, as measured by the Si~\\textsc{iv} O~\\textsc{iv}~]~/~C~\\textsc{iv} ratio, and outflow strength in quasars, as obtained via the blueshift and asymmetry index (BAI) of the C~\\textsc{iv} emission line. Their results indicate highly significant super--solar metallicity ($Z/Z_\\odot \\geq 5$) for quasars with BAI$\\geq 0.7$. This results indicates that metallicity likely plays an important role in the formation and acceleration of quasar outflows as expected, for instance, if quasar outflows are predominantly line--driven.  As mentioned above, both IRAS~13224--3809 and 1H~0707--495 are characterised by extremely blueshifted C~\\textsc{iv} emission lines with almost no contribution at rest wavelength. Their UV spectra indicate that BAI$\\geq 0.9$ in both sources, as shown in Fig.~\\ref{BAI}. If the metallicity--BAI correlation of Wang et al. (2012) extends or saturates above their largest observed BAI ($\\sim 0.76$), one infers that IRAS~13224--3809 and 1H~0707--495 are characterised by $Z/Z_\\odot \\geq 8$. A strong indication for super--solar metallicity in both sources is consistent with the strong FeII lines in the optical spectra and was also inferred by Leighly (2004) via photoionisation modelling of the UV spectra.  \\begin{figure} \\begin{center} \\includegraphics[width=0.35\\textwidth,height=0.45\\textwidth,angle=-90]{CivRatio.ps} \\caption{The C~\\textsc{iv} emission line profile from the HST--STIS observation performed on June 1999 with the G~140L grating is shown in the observed frame. Data have been slightly rebinned for visual clarity. The vertical line shows the expected wavelength of the C~\\textsc{iv} emission line for a redshift $z=0.0658$.} \\label{BAI} \\end{center} \\end{figure}     A $\\sim 100\\s$ soft lag is detected, which is a direct prediction of the reflection modelling used for the source. With the many other lags now seen, this justifies the reflection spectrum approach. It is consistent with the spectral modelling which indicates that the bulk of the primary continuum emission source is only a few gravitational radii in size and distance from the black hole. The spin of the black hole is high and close to maximal. This may be the result of secular evolution dominating in Narrow-Line Seyfert 1 galaxies, as inferred by Orban de Xivry et al (2011)."
    },
    "1208/1208.2689_arXiv.txt": {
        "abstract": "We present the first definitive measurement of the absolute magnitude of RR Lyrae c-type variable stars (RRc) determined purely from statistical parallax.  We use a sample of 247 RRc selected from the All Sky Automated Survey (ASAS) for which high-quality light curves, photometry and proper motions are available.  We obtain high-resolution echelle spectra for these objects to determine radial velocities and abundances as part of the Carnegie RR Lyrae Survey (CARRS).  We find that $M_{V,\\rm RRc} = 0.52 \\pm 0.11$ at a mean metallicity of ${\\rm [Fe/H]} = -1.59$.  This is to be compared with previous estimates for RRab stars ($M_{V,\\rm RRab} = 0.75 \\pm 0.13$) and the only {\\it direct} measurement of an RRc absolute magnitude (RZ Cephei, $M_{V, RRc} = 0.27\\pm 0.17$). We find the bulk velocity of the halo to be $(W_\\pi, W_\\theta, W_z) = (10.9,34.9,7.2)\\kms$ in the radial, rotational and vertical directions with dispersions $(\\sigma_{W_\\pi}, \\sigma_{W_\\theta}, \\sigma_{W_z}) =  (154.7, 103.6, 93.8) \\kms$.  For the disk, we find $(W_\\pi, W_\\theta, W_z) = (8.5, 213.2, -22.1)\\kms$  with dispersions $(\\sigma_{W_\\pi}, \\sigma_{W_\\theta}, \\sigma_{W_z}) =  (63.5, 49.6, 51.3) \\kms$. Finally, we suggest that UCAC2 proper motion errors may be overestimated by about 25\\%. ",
        "introduction": "Determining distances by use of multiple methods has a long and distinguished history in astronomy from antiquity to the present day.  Aristarchus of Samos first determined the Moon-Earth distance from the lunar eclipse.  A century later, Hipparchus checked Aristarchus' values using the independent method of terrestrial parallax: the position of the lunar limb during solar eclipse as seen from Alexandria and Hellespont.  More recently, astronomers have demanded that multiple methods be employed for zeroing in on the precise parameters governing the currently observed and mysterious accelerated expansion of the universe \\citep{albrecht06}.  Pulsating variables have enjoyed a privileged role in the local volume.  Due to their characteristic light curves and relatively bright absolute magnitudes compared to the bulk of main sequence stars, they can easily be identified and readily measured in nearby galaxies.  With a local zeropoint for these systems, it is straightforward, modulo metallicity and reddening issues, to determine the distances to external systems.  Cepheids, and their more common, albeit fainter, relatives within the instability strip, the x (RRL) variables, have been two key elements in the historical endeavor to launch humanity out of the solar system and Milky Way and into the local cosmos.  Indeed the Cepheids currently serve as the anchor of the cosmological distance scale, having allowed the most precise measurement the local rate of expansion of the universe, $H_0$, to date (Freedman et al. 1994, Freedman et al. 2001).  In principle, once properly calibrated, all stars of known absolute magnitude should yield identical measurements of the distances to nearby galaxies.  This has not historically been the case.  In particular, the RRL $M_V$ calibration has varied by almost half a magnitude depending on the method adopted, and, therefore, consistency with the Cepheid distance scale as well as other distance metrics has been difficult to establish.  As a result, attempts to determine RRL absolute magnitudes were largely abandoned over the past decade with a few notable exceptions (e.g., Dambis et al. 2009).   However, recently there have been new efforts.   Benedict et al. (2011) used HST trigonometric parallaxes of 5 RRL stars (including 4 RRab and 1 RRc variable) to obtain an average $M_V = 0.45 \\pm 0.05$.  Klein et al. (2011) recently used mid-IR data from the WISE satellite to infer a mid-IR Period-Luminosity relation.  In this work, we present a third measurement of RRL absolute magnitudes using the method of Statistical Parallax (\\statpi).   The large number of RRL that have been discovered in the last decade allow us to make a fresh assault on this issue.  Historically, RRL star distance measurements from  \\statpi\\ have come in systematically shorter than other distance indicators, in particular Cepheids (Barnes \\& Hawley 1986, Hawley et al. 1986, Layden et al. 1996 (hereafter, L96),  Popowski \\& Gould 1998a, Popowski \\& Gould 1998b , Gould \\& Popowski 1998 (hereafter, collectively PG$^3$), Dambis et al. 2009).  It is not yet fully understood why either this method or these two classes of objects should yield different distances to the same galaxies.   Thanks to automated synoptic all-sky surveys like the All Sky Automated Survey (ASAS, Pojmanski 2002), the number of RRL stars that have reliable light curves has increased by a factor $\\sim$5 relative to the previous ``state-of-the-art\".  The ASAS program has identified approximately 2000 RRL stars with 300-500 epochs of photometry.  By obtaining high-resolution spectra for these targets, we can both measure the radial velocities and metallicities needed for \\statpi\\ and address outstanding issues of systematics.  The light curves allow us to accurately determine pulsation phases and permit the measurement of the radial velocity at a single fiducial phase at which the pulsation velocity equals the star's systemic velocity (see Kollmeier et al. 2009).  Traditionally, obtaining the radial velocity component for RRL was laborious, requiring multiple epochs of spectroscopic observation.  The determination of the phase-velocity relationship for RRabc variables (Liu 1991, Kollmeier et al. 2009, Preston et al. 2011) allows a far more efficient strategy for obtaining critical radial velocity information with which to compute \\statpi\\ as we discuss further in Section 3.  Historically, RRc variables have been either excluded from \\statpi\\ analyses or only approximately analyzed.  This is primarily due to two factors.  First, their hotter temperatures make it more challenging to determine abundances from low-resolution, low signal-to-noise ratio (SNR) spectra (see Layden 1994) and, as a result, these objects cannot be robustly classified by population (halo/disk) as required by modern \\statpi.  However, high-resolution echelle observations circumvent this issue and allow, for the first time, a definitive \\statpi\\ analysis from RRc variables alone.  Second, there are fewer RRcs relative to RRabs, and it is only now that samples are large enough to perform a robust, self-consistent, pure RRc \\statpi\\ analysis.  We analyze our full (RRab + RRc) sample in a future work (Kollmeier et al. in preparation) and restrict our attention here to our RRc sample.  In Section 2 we present a brief overview of \\statpi\\ to remind the reader of the basic principles of the technique.  In Section 3 we present our sample selection, observations, data reduction, and analysis methods.  In Section 4 we review our updated methodology for determining \\statpi, the results of which are discussed in Section 5 and compared to previous \\statpi\\ results in Section 6.  Finally, in Section 7 we discuss our results in light of recent and historical works on the absolute magnitude scale of RRL variables. ",
        "conclusions": "We have performed the first decisive analysis of the absolute magnitude for RRc variables via statistical parallax using the first data from the Carnegie RR Lyrae Survey (CARRS).  Our current measurements for RRc variables yield a 5\\% distance error which is similar to that obtained from modern techniques applied to RRab samples.  We find a velocity ellipsoid for our disk and halo population that is in good agreement with previous measurements.  Already, CARRS provides competitive distance accuracy to other surveys and techniques.  At the conclusion of CARRS, we anticipate a factor of $\\sim 4$ increase in the number of tracers and, consequently, 2\\% distance errors.  In future work, we will analyze this far larger database and have the statistical potency to divide our sample into finer metallicity and kinematic bins than can be done presently.  This will allow a precision measurement for comparison with other techniques with the hope of a ``unified\" RRL distance scale.  In the Gaia era, where space-based parallaxes will be available for many of these objects, our database of high-resolution spectra should provide useful complementary information for going beyond distances and gaining further understanding of RRL as astrophysical objects, rather than merely as test particles.   \\appendix \\begin{appendices}"
    },
    "1208/1208.4250_arXiv.txt": {
        "abstract": "We used extensive ground-based multisite and archival spectroscopy to derive observational constraints for a seismic modelling of the magnetic $\\beta$~Cep star V2052~Ophiuchi. The line-profile variability is dominated by a radial mode ($f_1$=$7.14846$ d$^{-1}$) and by rotational modulation ($P_{\\rm rot}$=$3.638833$ d). Two non-radial low-amplitude modes ($f_2$=$7.75603\\ \\rm{d}^{-1}$ and $f_3$=$6.82308\\ \\rm{d}^{-1}$) are also detected. The four periodicities that we found are the same as the ones discovered from a companion multisite photometric campaign (Handler et al.\\,\\citealp{handler12}) and known in the literature. Using the photometric constraints on the degrees $\\ell$ of the pulsation modes, we show that both $f_2$ and $f_3$ are prograde modes with $(\\ell,m)$=$(4,2)$ or $(4,3)$. These results allowed us to deduce ranges for the mass ($M \\in [8.2,9.6]\\ $M$_\\odot$) and central hydrogen abundance ($X_c \\in [0.25,0.32]$) of V2052~Oph, to identify the radial orders $n_1$=$1$, $n_2$=$-3$ and $n_3$=$-2$, and to derive an equatorial rotation velocity $v_{\\rm eq} \\in [71,75]$ km~s$^{-1}$. The model parameters are in full agreement with the effective temperature and surface gravity deduced from spectroscopy. Only models with no or mild core overshooting ($\\alpha_{\\rm ov} \\in [0,0.15]$ local pressure scale heights) can account for the observed properties. Such a low overshooting is opposite to our previous modelling results for the non-magnetic $\\beta$~Cep star $\\theta$~Oph having very similar parameters, except for a slower surface rotation rate. We discuss whether this result can be explained by the presence of a magnetic field in V2052~Oph that inhibits mixing in its interior. ",
        "introduction": "Stellar modelling of B-type main-sequence stars is benefiting greatly from seismic data of $\\beta\\,$Cep stars. Among the various classes of B-type pulsators (e.g. Chapter\\,2 of Aerts et al.\\,\\citealp{aerts10}) this is the only one with members having clearly identified values for the wavenumbers $(\\ell,m,n)$ for several of the detected oscillation modes, which is a prerequisite for successful seismic modelling.   Models based on standard input physics could satisfactorily explain seismic data of the $\\beta\\,$Cep stars V836\\,Cen (Aerts et al.\\,\\citealp{aerts03}; Dupret et al.\\,\\citealp{dupret04}), $\\beta\\,$CMa (Mazumdar et al.\\,\\citealp{mazumdar06}), $\\delta\\,$Ceti (Aerts et al.\\,\\citealp{aerts06}), and $\\theta\\,$Oph (Briquet et al.\\,\\citealp{briquet07}). However, models with the same input physics failed to explain the much richer seismic constraints assembled from huge multi-technique multisite campaigns organised during many months for the $\\beta\\,$Cep stars $\\nu\\,$Eri (De Ridder et al.\\,\\citealp{deridder04}; Pamyatnykh, Handler \\& Dziembowski\\,\\citealp{pamyatnykh04}; Ausseloos et al.\\,\\citealp{ausseloos04}) and 12\\,Lac (Handler et al.\\,\\citealp{handler06}; Dziembowski \\& Pamyatnykh\\,\\citealp{dziembowski08}; Desmet et al.\\,\\citealp{desmet09}). These campaigns led to more than twice the number of modes known previously in those two stars, among which there were one or two high-order $g$-modes, and revealed shortcomings in the excitation predictions for some modes.  A similar excitation problem occurred for $\\gamma\\,$Peg, a class member found to pulsate in several high-order $g$-modes and low-order $p$-modes from MOST space-based photometry (Handler et al.\\,\\citealp{handler09}; Walczak \\& Daszy\\'nska-Daszkiewicz\\,\\citealp{walczak10}). Zdravkov \\& Pamyatnykh\\,\\citet{zdravkov09} suggested that an increase in the opacities by 20\\% might solve the excitation problem of $\\gamma\\,$Peg. Another excitation problem was pointed out for one mode of the CoRoT target V1449 Aql (Aerts et al.\\,\\citealp{aerts11}). The recent detection and interpretation of $\\beta\\,$Cep-type modes in the O9V pulsator HD\\,46202 (Briquet et al.\\,\\citealp{briquet11}) with CoRoT observations showed that the excitation problem of massive pulsators is acute, since none of the detected modes is predicted to be excited for appropriate stellar models representing this star. In all these studies, the effects of rotation on the excitation and amplitude of modes were not or only partly treated. Non-adiabatic computations taking rotation into account of the excitation of modes are available (e.g. Lee\\,\\citealp{lee98}; Townsend\\,\\citealp{townsend05}) but those of the amplitude of modes are just starting to become available (Lee\\,\\citealp{lee12}). These calculations will be very useful to check whether a better agreement between observed and theoretically excited modes may be found.  An important result of these asteroseismic studies is the evidence for non-rigid interior rotation in some $\\beta$~Cep stars. It was shown that V836~Cen, $\\nu$~Eri and 12~Lac rotate more rapidly in their inner parts than at their surface. These works also allow us to test whether core overshooting has to be included in our stellar models. Overshooting represents here the amount of non-standard mixing processes (Maeder\\,\\citealp{maeder09}; Mathis\\,\\citealp{mathis10}). For most of the $\\beta$~Cep targets modelled so far, the conclusion was the need for core overshooting for a better agreement with the pulsational characteristics. The derived core overshooting parameter values are however small (around 0.10 local pressure scale heights for V836~Cen and HD\\,46202, and around 0.20 for $\\delta$~Ceti, $\\beta$~CMa, and 12~Lac) with the exception of $\\theta$~Oph for which the value was found to be around 0.40. For two cases, $\\nu$~Eri and V1449 Aql, the core overshooting parameter could be kept to zero.   Space-based observations revealed other pulsators than $\\gamma\\,$Peg with both $g$- and $p$-mode pulsations of SPB and $\\beta$~Cep types (Balona et al.\\,\\citealp{balona11}). For the CoRoT hybrid HD~50230, the observation of deviations from a uniform period spacing led to the first exploration of the regions adjacent to the core in a massive star (Degroote et al.\\,\\citealp{degroote10}). HD~43317 was discovered by CoRoT to be another hybrid star with a wealth of pulsational constraints but also showing the presence of chemical inhomogeneities at its stellar surface (P\\'apics et al.\\,\\citealp{papics12}). Before space-based asteroseismology, the simultaneous presence of both pulsation and rotational modulation was found from ground-based spectroscopy only in $\\beta$~Cep (Telting, Aerts \\& Mathias\\,\\citealp{telting97}), V2052~Oph (Neiner et al.\\,\\citealp{neiner03}), and $\\kappa$~Sco (Uytterhoeven et al.\\,\\citealp{uytterhoeven05}). Another important discovery by means of space-based CoRoT data is the first observation of stochastically excited gravito-inertial modes in a massive star. Such modes have recently been detected in the hot (O9.5-B0) pulsator HD\\,51452 (Neiner et al., submitted), and this type of excitation should probably be more widely considered for massive pulsators.  Among $\\beta$~Cep stars with an asteroseismic modelling, a magnetic field has been detected so far in $\\beta$~Cep (Henrichs et al.\\,\\citealp{henrichs00}; Shibahashi \\& Aerts\\,\\citealp{shibahashi00}) and in V1449 Aql (Hubrig et al.\\,\\citealp{hubrig11}; Aerts et al.\\,\\citealp{aerts11}), although the field of the latter is still a matter of debate (Shultz et al.\\,\\citealp{shultz12}). In order to understand the effect of a magnetic field on the seismic behaviour of $\\beta\\,$Cep stars, a multisite photometric and spectroscopic campaign was set up for the equatorial star V2052~Oph ($m_V$=$5.8$, spectral type B2IV/V), which is known to be a magnetic pulsator with a dominant radial mode of frequency 7.145 d$^{-1}$ and a rotation frequency of 0.275 d$^{-1}$ (Neiner et al.\\,\\citealp{neiner03}). The star is slightly enriched in He, revealed a mild N excess (Morel et al.\\,\\citealp{morel06}), and was considered to be suitable as a seismic target, in view of the identified dominant radial mode and the discovery of an additional low-amplitude non-radial mode by Neiner et al.\\,\\citet{neiner03}.  The multisite photometry of V2052~Oph is reported in a companion paper by Handler et al.\\,\\citet{handler12}. Here, we present the spectroscopic part of the campaign. After a description of our dataset in Sect.\\,\\ref{sect_data}, we discuss the atmospheric parameters and chemical composition of the star (Sect.\\,\\ref{sect_parameters}). Sect.\\,\\ref{sect_lpv} is devoted to our frequency analysis and mode identification. Besides our line-profile study, we make a comparison with stellar models (Sect.\\,\\ref{sect_modelling}). We end with a discussion about the inhibition of mixing by magnetism and conclusions in Sect.\\,\\ref{sect_discussion} and Sect.\\,\\ref{sect_conclusions}, respectively. ",
        "conclusions": "Our study was based on extensive ground-based, high-resolution, high-S/N, multisite spectroscopic measurements spread over many years. Similar efforts were already performed for two other $\\beta$~Cep stars, $\\nu$~Eri and 12~Lac, leading to around ten independent pulsation frequencies. For V2052~Oph, our campaign only led to the detection of three pulsation modes but, contrary to the two other cases, the rotation period is also found, which provides a strong additional constraint. The difficulty to detect more modes in V2052~Oph may be explained by lower amplitudes of the modes due to a higher rotation rate compared to the two other stars.  Comparing our identifications of the degrees of the modes with the ones coming from the photometry, we are confident that the dominant mode is radial while the two additional modes have $\\ell$=$4$. Moreover, combining the outcome of the FPF method with that of our basic modelling, we conclude that the non-radial modes are prograde $(\\ell,m)$=$(4,2)$ or $(4,3)$ components of two multiplets with $n$=$-2$ and $n$=$-3$. The surface equatorial rotational velocity was also deduced as $v_{\\rm eq}$=73$\\pm$2 km~s$^{-1}$.  The observation of three independent modes with constraints on their wavenumbers along with the rotation period was enough to deduce ranges for the mass ($M \\in [8.2,9.6]\\ $M$_\\odot$), central hydrogen abundance ($X_c \\in [0.25,0.32]$) and other parameters of V2052~Oph that are furthermore fully compatible with the spectroscopic $T_{\\rm eff}$ and $\\log g$, and with the inclination angle $i$ determined from the modelling of the magnetic field. More importantly, we showed that only models with no or mild core overshooting could satisfactorily reproduce our observational properties.   We made a comparison with $\\theta$~Oph because it is an asteroseismically modelled $\\beta$~Cep star having similar fundamental parameters as our object of interest. We would expect a larger core overshooting parameter value in the faster rotator but the contrary is found. It can be explained by the fact that V2052~Oph is a magnetic star while $\\theta$~Oph is not. Indeed, using two approximate criteria, we showed that the magnetic field present in V2052~Oph is strong enough to inhibit non-standard mixing processes in its interior. In order to check if the internal rotation of V2052~Oph is indeed uniform or not, we would need additional constraints such as those obtained from a high-precision light curve as given by the CoRoT and Kepler satellites. Indeed, space-based photometry has already revealed many pulsation frequencies, even in objects that were thought to be monoperiodic from the ground (Aerts et al.\\,\\citealp{aerts06}; Degroote et al.\\,\\citealp{degroote09})."
    },
    "1208/1208.3536_arXiv.txt": {
        "abstract": "To constrain the origin of the soft X-ray excess phenomenon seen in many active galactic nuclei, the intensity-correlated spectral analysis, developed by Noda et al. (2011b) for Markarian 509, was applied to wide-band (0.5--45 keV) Suzaku data of five representative objects with relatively weak reflection signature. They are the typical bare-nucleus type 1 Seyfert  Fairall 9, the bright and typical type 1.5 Seyfert MCG-2-58-22, 3C382 which is one of the X-ray brightest broad line radio galaxies, the typical Seyfert-like radio loud quasar 4C+74.26, and the X-ray brightest radio quiet quasar MR2251-178. In all of them, soft X-ray intensities in energies below 3 keV were tightly correlated with that in 3--10 keV, but with significant positive offsets. These offsets, when calculated in finer energy bands, define a stable soft component in 0.5--3 keV. In each object, this component successfully explained the soft excess above a power-law fit. These components were interpreted in several alternative ways, including a thermal Comptonization component which is independent of the dominant power-law emission. This interpretation, considered physically most reasonable, is discussed from a viewpoint of Multi-Zone Comptonization, which was proposed for the black hole binary Cygnus X-1 (Makishima et al. 2008). ",
        "introduction": "\\label{sec:intro}  In soft X-ray spectra of Active Galactic Nuclei (AGNs), a phenomenon called ``soft X-ray excess'' is often observed. Characterized by a similar and steep flux upturn towards lower energies, this feature is clearly noticed particularly in weakly-absorbed AGNs, such as type I Seyferts and Broad Line Radio Galaxies (BLRG). For years, the origin of this spectral structure has been unidentified, and many interpretations have been proposed. Its simplest explanation is blackbody radiation from an optically-thick accretion disk. However, the observed color temperature of the soft excess, typically $\\sim 0.2$ keV, is too high for such disks around black holes (BHs) of  $\\gtrsim 10^7~M_\\odot$ in mass, where $M_\\odot$ is the solar mass.   \\begin{table*}[t] \\caption{Information of the five AGNs to be studied.} \\label{all_tbl} \\begin{center} \\begin{tabular}{ccccccc} \\hline\\hline Object name  &Type & Obs. date & redshift &$N_{\\rm H}$(Gal)$^{*}$  & reported $R^{\\dagger}$ & previous Suzaku study  \\\\  \\hline Fairall 9    \t    & Sy1 \t&   2007 June 7\t& $0.047$& 0.031   \t& 0.52$^{+0.20}_{-0.18}$\t& Patrick et al. (2011)\\\\ &\t&   2010 May 19\t&\t\t&\t&1.55$^{+0.26}_{-0.24}$  & Patrick et al. (2011) \\\\ MCG-2-58-22     & Sy1.5\t&    2009 December 2&   $0.047$ &\t0.027 \t& \t$0.69\\pm0.05$\t& Rivers et al. (2011)\\\\ 3C382    \t\t  & BLRG  &    2007 April 27\t&$0.058$\t&  0.074\t& 0.15$\\pm$0.05 \t& Sambruna et al. (2011)\\\\ 4C+74.26        & RLQ\t  &    2007 October 25\t&$0.104$& \t0.119 \t& \t0.3--0.7\t& Larsson et al. (2008)\\\\ MR2251-178 \t   & RQQ &   2009 May 7& $0.064$\t& \t0.024 \t& \t$\\leq$0.2\t & Gofford et al. (2011)\\\\  \\hline\\hline    \\end{tabular} \\end{center}  {\\small \\footnotemark[$*$] Equivalent hydrogen column density of the Galactic line-of-sight absorption in  $10^{22}$ cm$^{-2}$.  \\\\ \\footnotemark[$\\dagger$] Reflection fraction defined by  $R = \\Omega /2 \\pi$, when $\\Omega$ is the reflection solid angle.}    \\end{table*}    Given this, various alternative interpretations were proposed. Among them, mainly three ideas have been promising. One is absorption by  a partially-covering and ionized absorber often incorporating velocity smearing effects (e.g., Middleton \\& Done 2004; Schurch \\& Done 2008; O'Neil et al. 2007), another is relativistically blurred reflection from an ionized disk which is in similar conditions as the first case (Zoghbi et al. 2010; Nardini et al. 2011), and the other is a thermal Comptonization component that is separated from that producing the dominant Power-Law (PL) component (Marshall et al. 2003). Since these interpretations often degenerate in spectral analysis (Cerruti et al. 2011), there have been no conclusions for the origin of the soft X-ray excess in AGNs. Recently, however, Makishima et al. (2008) and Yamada (2011) established a Multi-Zone Comptonization (MZC) view for Cygnus X-1 (hereafter Cyg X-1) through the analysis of 0.5--200 keV Suzaku data, and showed that the last interpretation among the three indeed explains the soft excess (Frontera et al. 2001) seen in this leading black hole binary (BHB). This is analogous to the third interpretation of the AGN soft excess as described above.  To examine the origin of the soft X-ray excess in AGNs utilizing the wide-band Suzaku capability, Noda et al. (2011b) chose the typical weakly-absorbed type 1 Seyfert galaxy Markarian 509 (Mrk 509 hereafter), and developed a method to study how its 0.5--3 keV intensity in the Suzaku data is correlated to that in 3--10 keV, and found that a significant positive offset remains in the softer band intensity when the correlation is extrapolated to lower counts. Because the method is based on Count-Count Correlation with Positive Offset, hereafter, it is called ``C3PO'' method for simplicity. Thus, utilizing time variations with the C3PO method, they extracted a stable soft X-ray component, and successfully explained it by a thermal  Comptonization which is independent of the dominant PL continuum. When this soft Comptonization component is included, the time-averaged 0.5--45 keV spectrum of Mrk 509 was reproduced in terms of a weakly absorbed single PL and its reflection. In addition, Noda et al. (2011b) discovered that the new soft component varied on a  timescale longer than 3 days, independently of the PL variation. This securely excluded a competing interpretation of the detected soft component in terms of some largely extended thermal emission from the host galaxy.    At the same time as Noda et al. (2011b), Mehdipour et al. (2011) analyzed multi-wavelength data of the same object, Mrk 509, which were  obtained in a large campaign including XMM-Newton, Hubble, and FUSE. They obtained the same conclusion, that the soft X-ray excess of this AGN is created as a thermal Comptonization component which is independent of the principal PL continuum. Importantly, the parameters of the thermal Comptonization they derived, including an electron temperature of $\\sim 0.2$ keV, and an optical depth of $\\sim 16$, agree with those of Noda et al. (2011b). Thus, through the two independent methods (timing analysis and  multi-wavelength spectral fitting), the soft excess phenomenon in Mrk 509 has been confirmed to arise as Comptonization by a warm corona, which is presumably different from (though possibly related to) the hotter corona producing the harder PL continuum. When combined with the Suzaku results on Cyg X-1 quoted above, this result strengthens the general analogy between BHBs and AGNs.  The next step is to examine whether or not the same phenomenon as found in Mrk 509 and Cyg X-1 is present in a larger number of AGNs, including Seyferts and objects of other types. For this purpose, we utilize Suzaku archival datasets of AGNs, because the wide-band simultaneous coverage with Suzaku, typically available over a 0.5--45 keV band, is essential in determining the underlying continuum and the reflection component for each AGN, and hence to unambiguously identify the soft excess signals. In fact, many AGNs have been observed with Suzaku since its launch in 2005. However, the overall sample is neither complete nor homogeneous. Therefore, we choose to conduct the present study via the following two steps. The first is to define those AGN types which are suitable to our purpose.  The second is to select, from the Suzaku archive, the best target that represents each of the selected classes.  In the present work, errors refer to $\\pm 1\\sigma$ confidence limits, except model-fitting parameters in XSPEC for which $90\\%$ error ranges are adopted. ",
        "conclusions": "To study the origin of soft excess phenomena, five AGNs with different types were selected from the Suzaku archive, and analyzed; Fairall 9 (Sy1), MCG-2-58-22 (Sy1.5), 3C382 (BLRG), 4C+74.26 (RLQ), and  MR2251-178 (RQQ), which have X-ray signals dominated by emission from their central engines. Applying the C3PO method developed in Noda et al. (2011b) and Churazov, Gilfanov, and Revnivtsev (2001), who detected an SSE in the bright Sy1 Mrk 509 and the typical BHB Cyg X-1, respectively, we succeeded in extracting an SSE in the 0.5--3 keV band of all the five AGNs. In addition, the application of the same technique to the 15--45 keV HXD-PIN data allowed us to detect the SHE from at least two of the five AGNs.  The highly linear CCPs we obtained rule out possibilities that the soft X-ray variations of our sample AGNs are caused by changes in any absorption, including in particular a partially covering warm absorber which is one of the three main interpretations explaining the soft excess phenomena of AGNs (section 1). This is because a partially covering absorber, if variable in its covering fraction, column density, or ionization degree, would cause complex spectral changes, so the CCP would not show such a linear distribution as in figure 2--6. Furthermore, in this case the soft X-ray count would vanish, without leaving a positive offset, when the PL intensity becomes zero. Thus, the SSEs must be more independent of the PL continuum.  The SSE spectra, extracted from the five AGNs via the C3PO analysis, were reproduced successfully by five spectral models; PL, \\texttt{diskbb}, \\texttt{apec}, \\texttt{kdblur * reflionx} and \\texttt{comptt}. From physical considerations, two of them, the PL and \\texttt{diskbb}, were ruled out. Combining the static analysis with the dynamical one, \\texttt{apec} was found to be unrealistic. Furthermore, as given in figure 12, the measured 0.5--3 keV luminosity of the SSE of our AGNs is too high for thermal emission from their host galaxies (which would be at most $\\lesssim 10^{43}$ erg s$^{-1}$; e.g., Fukazawa et al. 2006), and correlates positively with the AGN luminosity. Therefore, the SSE must be tightly connected to the AGN phenomenon, rather than to the host galaxies.  \\begin{figure*}[t] \\begin{center} \\FigureFile(110mm,110mm) {figure12_low.eps} \\end{center} \\caption {The 0.5--3 keV luminosity of the SSE component (table 3) of the five AGNs, plotted against their 3--10 keV luminosity. Those of Mrk 509 reported in Noda et al. (2011b) are also calculated and plotted. }  \\label{fig:lcs} \\end{figure*}    How about the relativistically-smeared and ionized reflection interpretation modeled by \\texttt{kdblur * reflionx}? When a fine-tuned geometry is actually realized to enable the ``light bending'' effect (e.g., Miniutti et al. 2007), the reflection would not necessarily have to follow variations of the primary emission, and may remain stable like the SSE. To examine whether or not this interpretation is appropriate, the broad energy coverage with Suzaku is particularly important, because the reflection signals from an ionized disk are expected to appear both at the softest and hardest spectral ends. When figure 10 is closely examined, we find that the model invoking {\\tt kdblur * reflionx + pexrav} tends to under-predict the $<3$ keV portion of the time-averaged data, and over-predict the time-averaged HXD data. This effect indeed made the fit to MR2251-178 unacceptable. Furthermore, at least in MCG-2-58-22 and 4C+74.26, the {\\tt kdblur * reflionx + pexrav} component determined through our static analysis was not able to explain the dynamically determined SSE+SHE signals. Thus, the relativistic reflection interpretation cannot explain the broad-band Suzaku data of three out of the 5 AGNs. In contrast, the thermal Comptonization interpretation, using {\\tt comptt}, can consistently explain, in all the 5 AGNs, both the dynamically derived SSE and the static soft excess. This agreement between the two independent methods significantly strengthens the determination of the soft excess signals (or the SSE) in our sample AGNs. In this case, the SHE can be explained as a separate component arising from a cold reflector that is most likely located at a considerable distance from the central engine.    Presuming that the SSE is produced as a thermal Comptonization component which is separate from (but related to) the dominant PL, the corona of each AGN is then considered to consist of multiple regions having different optical depths and/or temperatures, namely, different Compton y-parameters. This may be called Multi-Zone Comptonization (MZC) condition. Since the various types of AGNs which have been selected in the present paper all have X-ray signals dominated by non-jet continua, their central engines (where most of the gravitational energy is converted to radiation) are inferred to be in the MZC condition. This agrees with the previous results on Mrk 509  (Noda et al. 2011b; Mehdipour et al. 2011), and the leading BHB Cyg X-1 (e.g., Makishima et al. 2008, Yamada  2011). The lack of significant short-term ($<$ several hundred ks) variability in the SSE may be explained if the SSE-producing Compton corona is largely ($>$ several hundred gravitational radii) extended, or more likely, if the seed-photon flux (presumably from the disk) is stable on these time scales while the dominant PL variations are produced in those part of the corona with the largest y-parameter (Makishima et al. 2008). Overall, we suggest that the soft excess emission of some (if not all) AGNs is actually a part of its primary continuum, produced in grossly the same central engine as the PL component, but possibly at somewhat different locations considering the clear difference in their variation characteristics. In this interpretation, the SHE is considered to have a different origin from the SSE, and produced via reflection by distant cool materials as represented by \\texttt{pexrav}.  The present results have impacts not only on the central engine, but also on the determination and interpretation of various secondary spectral components, including disk reflection, iron lines, and warm absorbers. This is because the MZC condition affects the primary continuum shape, which was often  assumed conventionally as a single PL. For example, Noda et al. (2011a) found that  the iron K$\\alpha$-line width of the type I Seyfert,  MCG--6-30-15, decreases considerably, when including a hard  component, which varied independently of the dominant PL and made the primary continuum concave.   We thank all members of the Suzaku hardware and software teams and the Science Working Group. HN, HU, and SS are supported by Japan Society for the Promotion of Science (JSPS) Research Fellowship for Young Scientists. KM is supported by Grantin-Aid for Scientific Research (A) (23244024) from JSPS, and SY by the Special Postdoctoral Researchers Program in RIKEN."
    },
    "1208/1208.0912_arXiv.txt": {
        "abstract": "We report the results of Giant Metrewave Radio Telescope (GMRT) observations of H{\\sc i} absorption towards the FRII radio galaxy 3C321 (J1531+2404), which is associated with an active galaxy interacting with a companion. The absorption profile towards the radio core is well resolved and consists three components, of which the two prominent ones are red-shifted by 186 and 235 km s$^{-1}$ relative to the optical systemic velocity. The neutral hydrogen column density towards the core is estimated to be $N$(H{\\sc i})=9.23$\\times$10$^{21}$(${T}_{\\rm s}$/100)($f_{c}$/1.0) cm$^{-2}$, where ${T}_{\\rm s}$ and $f_c$ are the spin temperature and covering factor of the background source respectively. We also present radio continuum observations of the source with both the GMRT and the Very Large Array (VLA) in order to understand the properties of a plume of emission at an angle of $\\sim$30$^\\circ$ to the source axis. This feature appears to have a steep high-frequency spectrum.  The current hotspots and jet are active and seen in X-ray emission. The spectral ages of the lobes are $\\lapp$26 Myr. We discuss the possibility that the  plume could be relic emission due to an earlier cycle of activity. ",
        "introduction": "Studying the properties of the gaseous environments of radio galaxies and quasars  on different scales are important for understanding the interactions of the radio jets with the external environment and the evolution of these sources. Such studies could also provide useful insights towards understanding the triggering of radio activity, and examining consistency of these properties with the unified schemes for active galactic nuclei (AGN) (Pihlstr\\\"om, Vermeulen \\& Conway 2003; Gupta \\& Saikia 2006b).  A useful way of probing the cold neutral component of this gas is via 21-cm H{\\sc i} absorption towards radio sources, which range in size from the sub-galactic sized compact steep-spectrum (CSS) and gigahertz peaked-spectrum (GPS) sources, to the large radio galaxies and quasars which are up to a few Mpc in size. The CSS and GPS objects (O' Dea 1998) have been inferred to be young ($<$10$^5$ yr), while the larger sources could be older than $\\sim$10$^8$ yr (Jamrozy et al. 2008; Konar et al. 2008).  Several H{\\sc i} absorption line studies have shown that CSS and GPS objects tend to exhibit absorption lines more frequently than larger sources, with the lines being both blue- and red-shifted relative to the systemic velocity and exhibiting complex line profiles. The H{\\sc i} column densities also appear to be anticorrelated with the source sizes (Pihlstr\\\"om, Conway \\& Vermeulen 2003; Gupta et al. 2006), and are broadly consistent with the unified scheme for radio galaxies and quasars. The relationship between the H{\\sc i} column density and core prominence, which is being used as a statistical indicator of source orientation, appears to be consistent with the H{\\sc i} gas being distributed in a circumnuclear disk on a scale smaller than the size of the compact radio sources (Gupta \\& Saikia 2006b). Evidence of absorption arising from a circumnuclear disk-like structure has also been inferred from higher-resolution, Very Long Baseline Interferometric (VLBI)-scale spectroscopic observations in several sources such as the CSS object J0119+3210 (4C+31.04) and the cores of a few larger sources, Cyg\\,A, NGC\\,4261 and Hydra\\,A (e.g. Conway \\& Blanco 1995; Taylor 1996; Conway 1999; van Langevelde et al. 2000). A comparison of the H{\\sc i} absorption properties towards the cores of larger sources with those of the CSS and GPS objects might provide insights towards understanding the evolution of the gaseous properties as the source ages. Such a study indicates that the detection of H{\\sc i} in absorption towards the cores of larger sources is significantly smaller than for CSS and GPS objects suggesting an evolution in the gaseous content of the host galaxies with source age (Chandola et al., in preparation). There have also been suggestions that the torus/disk may be different in FR\\,I and FR\\,II sources, which is likely to be reflected in the H{\\sc i} absorption properties of these two classes of sources (e.g. Morganti et al. 2001). Given the small number of detections at present, the difference in detection rate between FR\\,I and FR\\,II sources does not appear to be significant (Chandola et al., in preparation).  An important aspect in our understanding of AGN is the episodic nature of their nuclear activity, and the physical processes that might be governing it (Saikia \\& Jamrozy 2009 for a review). There appears to be a trend for a high rate of detection of H{\\sc i} absorption in sources with evidence of rejuvenated activity. These include the giant radio galaxy 3C236 which also exhibits evidence of star formation (Conway \\& Schilizzi 2000), the giant radio galaxy J1247+6723 with a GPS core (Saikia, Gupta \\& Konar  2007), the misaligned DDRG 3C293 (Beswick et al. 2004; Emonts et al. 2005), the well-studied southern radio galaxy Centaurus A (Sarma, Troland \\& Rupen 2002; Morganti et al. 2008), 4C~29.30 (Chandola, Saikia \\& Gupta 2010) and CTA~21 (Salter et al. 2010). The well-known FRII radio galaxy Cygnus A, which has been shown to have two cycles of radio activity from radio and X-ray observations (Steenbrugge, Blundell \\& Duffy 2008; Steenbrugge, Heywood \\& Blundell 2010), also exhibits nuclear H{\\sc i} absorption (Conway 1999).  In this paper we present H{\\sc i} observations with the Giant Metrewave Radio Telescope (GMRT) to localise and put constraints on the size of the H{\\sc i} absorber in the FRII radio galaxy 3C321, which has been reported earlier from Arecibo observations (Mirabel 1990). We also present radio continuum observations with the GMRT and the Very Large Array (VLA) of the diffuse plume of emission to explore whether this feature might be relic emission from an earlier cycle of activity. ",
        "conclusions": "The GMRT H{\\sc i} observations show that the absorption occurs towards the core component, with also a suggestion of absorption towards the knot in the jet 6 kpc north-west of the nucleus. This would indicate the size of the absorber to be $\\gapp$ 6 kpc if it is a single absorber. The absorption profile towards the core has multiple components with a total column density of $N$(H{\\sc i})=9.23$\\times$10$^{21}$(${T}_{\\rm s}$/100)($f_c$/1.0) cm$^{-2}$. The column density towards the knot is $N$(H{\\sc i})=7.54$\\times$10$^{21}$(${T}_{\\rm s}$/100)($f_c$/1.0) cm$^{-2}$. No absorption is seen towards the hot-spots.  Radio continuum observations of lobes with similar resolution show  that their spectra are straight up to 5 GHz, suggesting spectral ages of $\\lapp$26 Myr for the lobes. The detection of X-rays from the hot-spots suggests that these are being continuously fed by the jets, consistent with their spectral ages. The full extent of the plume is not seen in our 614-MHz GMRT and also barely seen in the TGSS image at 148 MHz. However, better quality data are required to determine the spectrum of the plume reliably over a large frequency range to try and examine a possible break in the spectrum and explore whether the plume might be relic emission from an earlier cycle of activity."
    },
    "1208/1208.0123_arXiv.txt": {
        "abstract": "We report on a comprehensive study of the phase structure of cold, dilute nuclear matter featuring a $\\SD$ condensate at non-zero isospin asymmetry, within wide ranges of temperatures and densities. We find a rich phase diagram comprising three superfluid phases, namely a Larkin-Ovchinnikov-Fulde-Ferrell phase, the ordinary BCS phase, and a heterogeneous, phase-separated BCS phase, with associated crossovers from the latter two phases to a homogeneous or phase-separated Bose-Einstein condensate of deuterons.  The phase diagram contains two tricritical points (one a Lifshitz point), which may degenerate into a single tetracritical point for some degree of isospin asymmetry. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0315_arXiv.txt": {
        "abstract": "In order to test a recent hypothesis that the dispersion in the Schmidt-Kennicutt law arises from variations in the evolutionary stage of star forming molecular clouds, we compared molecular gas and recent star formation in an early-phase merger galaxy pair, Taffy I (UGC\\ 12915/UGC\\ 12914, VV\\ 254) which went through a direct collision 20 Myr ago and whose star forming regions are expected to have similar ages.  Narrow-band Pa$\\alpha$ image is obtained using the ANIR near-infrared camera on the mini-TAO 1m telescope.  The image enables us to derive accurate star formation rates within the galaxy directly.  The total star formation rate, $22.2\\ M_\\odot \\mathrm{yr^{-1}}$, was found to be much higher than previous estimates.  Ages of individual star forming blobs estimated from equivalent widths indicate that most star forming regions are $\\sim$ 7 Myr old, except for a giant HII region at the bridge which is much younger.  Comparison between star formation rates and molecular gas masses for the regions with the same age exhibits a surprisingly tight correlation, a slope of unity, and star formation efficiencies comparable to those of starburst galaxies.  These results suggest that Taffy I has just evolved into a starburst system after the collision, and the star forming sites are at a similar stage in their evolution from natal molecular clouds except for the bridge region.  The tight Schmidt-Kennicutt law supports the scenario that dispersion in the star formation law is in large part due to differences in evolutionary stage of star forming regions. ",
        "introduction": "Understanding the nature and physics which underly the relation between star formation rate (SFR) and gas density is crucial for understanding galaxy evolution and star formation history of the universe.  The most thoroughly investigated expression of this relation is the Schmidt-Kennicutt (SK) law \\citep{kennicutt98}, the empirical power-law relation between surface densities of SFR ($\\Sigma_\\mathrm{SFR}$) and  molecular gas ($\\Sigma_\\mathrm{H2}$), written as \\begin{equation} \\Sigma_\\mathrm{SFR} \\propto \\Sigma_\\mathrm{H2}^N. \\end{equation} The power law index $N$ is typically found to be $N>1$ for the commonly used molecular gas tracers $^{12}\\mathrm{CO}(J=1-0)$ \\citep{heyer04, komugi05} and $^{12}\\mathrm{CO}(J=2-1)$ \\citep{schuster07}, but $N\\sim 1.0$ when dense gas tracers such as high CO transitions and HCN are used to trace $\\Sigma_\\mathrm{H2}$ \\citep{gao04, komugi07, bigiel08, muraoka09, iono09}.  This is typically interpreted as the dense gas tracing individual units of star formation \\citep{wu05}, which forms stars at a constant efficiency.   A less emphasized aspect of the SK law is its dispersion around the relation.  Much of the dispersion can be attributed to uncertainties in calibrating $\\Sigma_\\mathrm{SFR}$ and $\\Sigma_\\mathrm{H2}$, as frequently used massive star formation tracers such as H$\\alpha$ are strongly affected by dust extinction \\citep{komugi07}, and molecular gas mass depends on the assumed CO-to-$\\mathrm{H_2}$ conversion factor, $\\mathrm{X_{CO}}$, which can vary significantly with environment \\citep[e.g.,][]{arimoto96, israel97, komugi11}. SK laws studied over a scale of several kilo-parsecs \\citep{kennicutt98, komugi05, schuster07, onodera10} typically scatter at least factor of $\\sim 2$. Recent studies \\citep{kennicutt07, schruba10, onodera10, liu11} have found that the dispersion in the SK law increases as smaller regions are sampled within the galaxy, which are unlikely to be explained by observational uncertainties.  The significant dispersion is attributed to differences in the evolutionary stage of the molecular clouds \\citep{onodera10}, which may be a large scale manifestation of molecular cloud types proposed by \\citet{kawamura09}.  An interesting test of this hypothesis is to see whether the dispersion of the SK law decreases significantly when star forming molecular clouds of comparable age are sampled.  An interacting galaxy system is an ideal target for this purpose.  Star formation is expected to be triggered by galaxy collision \\citep{barnes96, saitoh09}, and indeed interacting galaxies such as the Antennae \\citep{whitmore99, wilson00, ueda12} and IIZw 096 \\citep{inami10} are known to host exceptionally young stellar clusters in the overlapping region and tidal tails.  If the interaction occurs over a timescale that is less than a typical timescale of star formation or molecular cloud evolution, and recent enough so that there has been time for only one generation of star formation, then the galactic collision-induced star formation in various regions of the system are expected to be at a comparable evolutionary stage.  The interacting galaxy pair Taffy I, is a rare laboratory which meets these requirements.   \\subsection{Taffy I} Taffy I (see figure 1) is an interacting system consisting of two galaxies (UGC12914/15) and the ``bridge'', an extended shock-induced synchrotron component connecting the two galaxies \\citep[hereafter C93]{condon93}. The distance is estimated to be 61 Mpc, from the systemic HI velocity of $4600\\ \\mathrm{km s^{-1}}$ (C93) and Hubble constant of $75\\ \\mathrm{km s^{-1} Mpc^{-1}}$. The steepening of spectral index at the bridge and the simulated trajectory of the galaxies indicate a face-on collision which occurred only 20 Myr ago (C93). The galaxies are separating at a velocity of $\\sim 450\\ \\mathrm{km\\ s^{-1}}$ for the assumed Hubble constant (C93), and the relative inclination of the disks is $12^\\circ$ \\citep{giovanelli86}.  For a star forming disk of 7 kpc diameter (the separation between star forming regions E and G in UGC12915; see following sections and figure 1), the duration of the disk-disk collision is estimated to be $\\sim 3$ Myr. This is much less than the typical timescale of star formation in galactic disks \\citep[several $\\times$ 10 Myr;][]{egusa04, egusa09} or the timescale of giant molecular clouds to evolve \\citep[6-13 Myr;][]{kawamura09}.  Thus, the collision was instantaneous in terms of the star formation history of this system, and recent enough so that the triggered star formation would not have evolved for more than a generation.  Star formation is evident from H$\\alpha$ and mid infrared observations, in both galactic disks and notably in the bridge region, where a giant HII region has been found \\citep{bushouse87}.  The SFR of this system has not been well measured, however, because commonly used tracers fail to trace massive star formation reliable in this galaxy.  H$\\alpha$ observation \\citep{bushouse87} indicates a total SFR of $\\sim 1.4\\ M_\\odot \\mathrm{yr^{-1}}$ (see section 3.2) .  However, the center of UGC\\ 12915 was virtually undetected in H$\\alpha$, indicating severe dust extinction. \\citet{gao03} used the 20cm continuum to derive the star formation rate, but cautioned that the 20cm is contaminated largely in the bridge by the shock-induced synchrotron emission (C93) which is unrelated to the massive star formation.  Seven micron PAH features have been observed by \\citet{jarrett99}, which returned a low global star formation rate of $ <\\ 6\\ M_\\odot \\mathrm{yr^{-1}}$ if we assume a conversion factor to SFR by \\citet{roussel01}, but using a common conversion factor of mid-infrared PAH emission to SFRs is debatable in Taffy I, as the small dust particles traced at these wavelengths could have been destroyed by the strong shock induced by the large scale galactic collision \\citep{braine03}.  The total infrared luminosity of L(IR) $ = 7\\times 10^{10}\\ L_\\odot$ corresponds to a SFR of $12.1\\ M_\\odot \\mathrm{yr^{-1}}$ using the conversion factor by \\citet{kennicutt98}.  This estimate is larger compared to those using other tracers, but the conversion factor assumes continuous and constant star formation over 10-100 Myr, which is questionable in these colliding galaxies.  Obtaining a more reliable estimation of the SFR within Taffy I and its age, requires an unbiased measure of massive star formation.  The Pa$\\alpha$ emission line ($\\lambda = 1.8751\\ \\mathrm{\\mu m}$) is suited for this goal, as the magnitude of its dust extinction is only 1/6 of that of H$\\alpha$, and can trace massive star formation more directly than infrared dust emission.  Therefore, we have carried out a narrow-band imaging observation of the redshifted Pa$\\alpha$ emission line. ",
        "conclusions": "\\subsection{Age} Equivalent widths of the Pa$\\alpha$ line are derived by dividing the Pa$\\alpha$ image with the $1.91 \\mathrm{\\mu m}$ continuum. Approximate ages of the blobs are estimated from the equivalent widths using figure 5 of \\citet{diaz08} which is based on Starburst99 \\citep{leitherer99} with an instantaneous burst, and shown in Table 1. The signal to noise ratio (SNR) of the Pa$\\alpha$ blobs within the measured aperture ranged from SNR $\\sim$7 to over 20, except for Region H in the faint tail of UGC12915, for which the SNR was 3.3.  The ages of all blobs except for region B (UGC\\ 12914 center) and region D (the bridge) are $\\sim 7$ Myr.  Region D is significantly younger, less than 3.5 Myr.  Region B is somewhat older, which is likely a contribution from an older population of stars as this is the center of the galaxy. In any case, the estimated ages of the star forming blobs are less than 20 Myr, the estimated time since the collision of the two galaxies. The ages are surprisingly uniform considering the time since collision, but this may be explained by the Pa$\\alpha$ observation sampling the brightest regions in the system which are forming stars most actively at this point.   \\subsection{Star Formation Efficiency and Schmidt-Kennicutt Law} The total molecular gas mass of this system is estimated to be $M(\\mathrm{H_2}) = 1.0\\times 10^{10}\\ \\mathrm{M_\\odot}$ (\\citet{braine03}; Kaneko et al. submitted) using a CO-to-$\\mathrm{H_2}$ conversion factor of $5.6\\times 10^{19}\\ \\mathrm{cm^{-2} [K\\ km\\ s^{-1}]^{-1}}$ \\citep{zhu07}.  From the global SFR estimate of $22.2\\ M_\\odot \\mathrm{yr^{-1}}$, the average star formation efficiency (SFE=SFR/$M(\\mathrm{H_2})$) in Taffy I is $2.2 \\times 10^{-9} \\mathrm{yr^{-1}}$, corresponding to a gas consumption timescale $\\tau_\\mathrm{gas}$ of $4.5\\times 10^8$ yr. This gas consumption timescale is consistent with typical starburst galaxies \\citep{kennicutt98}, and qualifies Taffy I as a starburst galaxy in contrast to indications from previous studies \\citep{gao03}.  In order to derive the molecular gas mass and SFE of individual blobs, we use interferometric observations of the $\\mathrm{^{12}CO}(J=1-0)$ line \\citep{iono05}.  The angular resolution of the CO data is $7^{\\prime \\prime}.2 \\times 5^{\\prime \\prime}.1$, sufficient to match the aperture ($12^{\\prime \\prime}.6$) used to determine the SFRs and the equivalent widths of the Pa$\\alpha$ emission. For the conversion factor, we use both the average $\\mathrm{X_{CO}}$ and the values derived individually for UGC12914, UGC12915 and the bridge by \\citet{zhu07} using a one-zone LVG analysis.  The derived gas masses are shown in Table \\ref{blobs} for both cases. The SFEs of the blobs are surprisingly uniform, with $\\mathrm{SFE}=(7.2\\pm 0.9) \\times 10^{-9} \\mathrm{yr^{-1}}$, corresponding to  $\\tau_\\mathrm{gas}=(1.4 \\pm 0.2) \\times 10^8$ yr, when the average conversion factor is used. When different conversion factors are used for the galaxies and the bridge, the bridge has an exceptionally high SFE and short gas consumption timescale, with $\\mathrm{SFE}=(1.4 \\pm 0.1) \\times 10^{-8} \\mathrm{yr^{-1}}$ and $\\tau_\\mathrm{gas} =(7.1\\pm 0.7) \\times 10^7$ yr, whereas all the other regions have nearly constant SFE and $\\tau_\\mathrm{gas}$, of $\\mathrm{SFE}=(4.9\\pm 0.9) \\times 10^{-9} \\mathrm{yr^{-1}}$ and $\\tau_\\mathrm{gas}=(2.0 \\pm 0.4) \\times 10^8$ yr.  These values are all typical of circumnuclear starburst regions in nearby galaxies \\citep{kennicutt98}.    The variation in the SFEs is also apparent from the comparison between molecular gas mass and SFR, shown in Figure \\ref{schmidt}. Since the projected areas of the apertures are same for all the blobs, Figure \\ref{schmidt} can also be regarded as the SK law.  For the constant conversion factor, a best-fit to all the blobs results in \\begin{equation} \\log \\Sigma_\\mathrm{SFR}=(0.99 \\pm 0.08)\\log \\Sigma_\\mathrm{H2} - (8.13 \\pm 0.12) \\end{equation} and when using different conversion factors, all blobs excluding region D (bridge) results in \\begin{equation} \\log \\Sigma_\\mathrm{SFR}=(0.95 \\pm 0.14)\\log \\Sigma_\\mathrm{H2} - (8.23 \\pm 0.24) \\end{equation} The observed power law index of $N=1.0$ is consistent with those typically observed using dense gas tracers, but not with with the typically observed $N>1$ using $\\mathrm{^{12}CO}(J=1-0)$.   \\subsection{Schmidt-Kennicutt Law and Dispersion} The dispersion $\\sigma$ of the correlation from the best fit is 0.058 dex (14\\%) in case of using a constant conversion factor, or 0.096 dex (25\\%) when using different conversion factors and excluding region D. The $\\sigma$ are comparable to or smaller than uncertainties in the SFR (25\\%), so uncertainties in the measurement alone can account for the dispersion for all regions with comparable ages, and only the exceptionally young region D is offset significantly from this correlation.  The dispersion can be compared to those in normal galaxies which contain star forming regions at a variety of evolutionary stages.  Although we use an aperture with a projected scale of 3.7 kpc, this is limited by the CO resolution.  Inspection of the high resolution Pa$\\alpha$ image reveals that even with this large aperture size, only one bright star forming complex is enclosed in each of the apertures. The effect of averaging over star forming regions which work to decrease $\\sigma$ \\citep{liu11} would not be prominent in our aperture, so we compare the $\\sigma$ in the literature at a common scale of $\\sim 700$ pc, the typical size of Pa$\\alpha$ blobs in Taffy I.  \\citet{liu11} derive $\\sigma = 0.50$ and $0.27$ for M51 and NGC3521, respectively, at 700 pc resolution.  Similarly, we derive a dispersion of $\\sigma = 0.43$ from the M33 dataset published by \\cite{onodera10} at 500 pc resolution, and $\\sigma = 0.32$ at 1 kpc resolution.  Given that our star forming blobs are distributed widely both within and between galaxies, the measured dispersion in Taffy I is unexpectedly small.  \\subsection{Star Formation History and Evolution of Taffy I} A stellar mass of a galaxy reflects its integrated star formation history, and therefore the ratio of current star formation rate to total stellar mass (the \\textit{specific star formation rate}; sSFR) indicates the intensity of current star formation relative to its average in the past.  The $K$ band magnitudes of UGC12914 ($9.32\\ \\mathrm{mag.}$) and UGC12915 ($9.83\\ \\mathrm{mag.}$; \\cite{jarrett99}) can be used to estimate their rotation velocities and hence their stellar mass based on the Tully-Fischer relation \\citep{bell01}, which give stellar mass estimates of $M_* = 7\\times 10^{10}\\ M_\\odot$ and $M_* = 5\\times 10^{10}\\ M_\\odot$ for UGC12914 and UGC12915, respectively.  Using the SFR of $22.2\\ M_\\odot \\mathrm{yr^{-1}}$, the sSFR of Taffy I is $1.9\\times 10^{-10}\\ \\mathrm{yr^{-1}}$.  This is five times larger than the average of local ($0 \\le z \\le 0.2$) galaxies with similar stellar mass ($\\sim 4\\times 10^{-11}\\ \\mathrm{yr^{-1}}$) given by \\citet{oliver10}, indicating that Taffy I is now forming stars much more actively than its average past.  The ages of the blobs estimated in the previous section also indicate that the star forming regions studied here are experiencing their first burst of star formation. Apparently, Taffy I has just evolved into a starburst system as a result of the collision which triggered massive star formation.       A scenario of star formation in Taffy I should explain the matched ages of the all regions except region D, and also the apparent delay of star formation in region D.   \\citet{braine03} and \\citet{zhu07} suggest that the molecular gas in the bridge was ionized in the collision but cooled promptly, recombined and returned to the molecular state in several to 10 Myr \\citep{bergin04}, delaying star formation by that time.  This scenario requires the cooling time of the gas which created region D to be longer than the timescale of collision ($\\sim$ 2 Myr or less: see section 1.1).  The cooling time of molecular gas at $10^{1-2} \\ \\mathrm{cm^{-3}}$ and $10^4$ K is $10^{4-6}$ years, so the colliding gas may not have had enough time to be heated to delay molecular gas formation compared to other regions. Another possible scenario is that the molecular gas in region D was simply stretched as a result of tidal force, delaying aggregation of gas.  Assuming that UGC12915 is the dominant source of tidal force and has $M_* = 5\\times 10^{10}\\ M_\\odot$, and using molecular mass of $1.2 \\times 10^8 \\ \\mathrm{M_\\odot}$ and diameter of 1 kpc for Region D, the Roche limit is reached at 5 kpc.  The projected distance of region D to UGC12915 is $\\sim$ 4 kpc, so the tidal force may be strongly affecting region D.  It is plausible that only recently the gas has been able to collapse locally within region D.  Apart from region D, the blobs can be regarded to be in a similar stage in their evolution from molecular gas to stars.  The SK law is consistent in this respect, as all the blobs except region D follow an exceptionally tight linear law, which is predicted if our working assumption that variations in the evolutionary stage of molecular clouds are responsible for the dispersion is correct.  Since the observed slope of $N = 1.0$ is typically observed with dense gas tracers, we speculate that the blobs in the same evolutionary stage have the same dense gas fraction, giving rise to a linear slope even when observed with $^{12}\\mathrm{CO}(J=1-0)$ that traces the total molecular mass.  This can be tested by high resolution observations of dense molecular gas tracers, which should show that the dense gas fraction in these blobs are similar."
    },
    "1208/1208.6369_arXiv.txt": {
        "abstract": "Quantum interference phenomena manifests itself in several ways in the polarized solar spectrum formed due to coherent scattering processes. One such effect arises due to interference between the fine structure $(J)$ states giving rise to multiplets. Another effect is that which arises due to interference between the hyperfine structure $(F)$ states. We extend the redistribution matrix derived for the $J$-state interference to the case of $F$-state interference. We then incorporate it into the polarized radiative transfer equation and solve it for isothermal constant property slab atmospheres. The relevant transfer equation is solved using a polarized approximate lambda iteration (PALI) technique based on operator perturbation. An alternative method derived from the Neumann series expansion is also proposed and is found to be relatively more efficient than the PALI method. The effects of PRD and the $F$-state interference on the shapes of the linearly polarized Stokes profiles are discussed. The emergent Stokes profiles are computed for hypothetical line transitions arising due to hyperfine structure splitting (HFS) of the upper $J=3/2$ and lower $J=1/2$ levels of a two-level atom model with nuclear spin $I_s=3/2$. We confine our attention to the non-magnetic scattering in the collisionless regime. ",
        "introduction": "\\label{intro} The linearly polarized solar spectrum is produced by coherent scattering processes taking place in the solar atmosphere. This so called `second solar spectrum' is highly structured and reveals various physical processes responsible to generate the polarized signals in the spectrum. Quantum interference is one such physical process whose importance has been highlighted in the second solar spectrum studies \\citep[see][]{s80,jos94,s97}. The coherent superposition of the fine structure states leads to the $J$-state interference, whereas the $F$-state interference arises due to superposition of the hyperfine structure states (see Figure~{\\ref{level-diag}}). The $J$-state interference theory for the case of frequency coherent scattering was developed by \\citet[][]{s80,jos94,s97}. This theory was extended to include partial frequency redistribution (PRD) in line scattering, by \\citet[][hereafter called P1]{smi11a}. The $J$-state PRD matrix derived in P1 is used in the polarized line transfer equation in \\citet[][herafter P2]{smi11b}. An alternative scattering theory of $J$-state interference based on metalevel approach was developed by \\citet[][]{landietal97}, which also includes the $F$-state interference effects. All the papers mentioned so far are applicable to the case of colisionless regime.  Second solar spectrum contains several lines which have signatures of $F$-state interference. Examples of these lines are Na {\\sc i} D$_2$ at 5890\\,\\AA, Ba {\\sc ii} D$_2$ at 4554\\,\\AA , Mn {\\sc i} 8741\\ \\AA, Sc {\\sc ii} 4247\\,\\AA\\ etc. In this paper we are concerned with the line formation studies involving $F$-state interference process and PRD. The $F$-state redistribution matrix derived in this paper can be used for modeling the non-magnetic quiet region observations of hyperfine structure splitting (HFS) in the lines mentioned above.   The $F$-state interference theory applicable to the frequency coherent scattering was developed by \\citet[][]{s97}. This theory, along with PRD, was applied by \\citet[][]{fluri03} and \\citet[][]{holz05} in the polarized line transfer computations. In \\citet[][]{ll04} the theory of $F$-state interference was developed under the approximation of complete frequency redistribution (CRD). The theory of $F$-state interference in a magnetic field for multi-term atoms in the collisionless regime and under the approximation of CRD is presented in \\citet[][]{manso05}.  In the present paper we extend the $J$-state interference theory presented in P1 to the case of $F$-state interference. The $F$-state redistribution matrix is derived here for the non-magnetic case and in the collisionless regime. The reason for considering the non-magnetic case in this paper, is that the formulation of P1 was confined to the linear Zeeman regime of field strengths (the spacing between the Zeeman $m$-states being smaller than the spacing between the fine structure states). In the present context, the hyperfine splitting becomes comparable to the Zeeman splitting even for weak magnetic fields, and we quickly enter the Paschen-Back regime of field strengths (level crossing of the $m$-states belonging to different $F$-states), in which the formulation presented in P1 is not valid. Since the Paschen-Back effect is outside the scope of the present paper, our treatment here is limited to the non-magnetic case, but the extension to the Paschen-Back regime is planned to be pursued in future work. We further assume that the lower level is unpolarized and infinitely sharp. While this assumption is made for the sake of mathematical simplicity, it is physically justified for the long-lived ground states, which are correspondingly more vulnerable to collisional depolarization.  Following P2, this PRD matrix is incorporated into the polarized line transfer equation, and solved using an operator perturbation method. We also propose a new method to solve the $F$-state interference problem. It is called the scattering expansion method (SEM) and is described in \\citet[][]{hfetal09,sam11}. Recently it has been applied to a variety of problems \\citep[see][]{sow12,sup12}. We compare the operator perturbation method and the SEM by applying them to the problem at hand.  In Section~{\\ref{theory}} we derive the PRD matrix for $F$-state interference and incorporate it into the line transfer equation. In Section~{\\ref{num-method}} we describe the numerical methods used to solve the transfer equation. Results are presented in Section~{\\ref{results}}. Section~{\\ref{conclusions}} is devoted to the concluding remarks. \\begin{figure} \\centering \\includegraphics[width=4.0cm,height=3.0cm]{f1.eps} \\caption{\\footnotesize Level diagram representing the HFS in a two-level atom model.} \\label{level-diag} \\end{figure} ",
        "conclusions": "\\label{conclusions} In the present paper we have extended the $J$-state interference formulation discussed in P1 and P2 to the case of $F$-state interference. The treatment is restricted to the collisionless and non-magnetic regime. The decomposition technique presented in \\citet{hf07} is applied to the $F$-state interference problem. It helps to incorporate the $F$-state interference redistribution matrix into the reduced form of the line radiative transfer equation. The transfer equation is solved using the traditional PALI and the scattering expansion method by suitably adapting them to handle the $F$-state interference problem. The SEM is found to be more efficient and faster than the PALI method.  The Stokes profiles computed by taking account of HFS are similar to the profiles of a single line arising from a two-level atom model without HFS. The HFS causes a depolarization of $Q/I$ in the line core irrespective of whether the $F$-state interference is taken into account or not. Like the $J$-state interference, the $F$-state interference affects mainly the line wing PRD peaks. We also show that when interference effects are neglected, the principle of spectroscopic stability is violated in both single scattered and multiple scattered profiles. Using the fractional polarization $Q/I$ in the $90^\\circ$ single scattering case, we can numerically estimate the wavelength dependent polarizability factor $W_2(\\lambda)$. The $W_2(\\lambda)$ so computed can then be used in the transfer equation to compare with our exact redistribution matrix approach. The two approaches are found to give identical emergent Stokes profiles."
    },
    "1208/1208.1852_arXiv.txt": {
        "abstract": "Superfluid hydrodynamics affects the spin-evolution of mature neutron stars, and may be key to explaining timing irregularities such as pulsar glitches. However, most models for this phenomenon exclude the global instability required to trigger the event. In this paper we discuss a mechanism that may fill this gap. We establish that small scale inertial r-modes become unstable in a superfluid neutron star that exhibits a rotational lag, expected to build up due to vortex pinning as the star spins down. Somewhat counterintuitively, this instability arises due to the (under normal circumstances dissipative) vortex-mediated mutual friction. We explore the nature of the superfluid instability for a simple incompressible model, allowing for entrainment coupling between the two fluid components. Our results recover a previously discussed dynamical instability in systems where the two components are strongly coupled. In addition, we demonstrate for the first time that the system is secularly unstable (with a growth time that scales with the mutual friction) throughout much of parameter space. Interestingly, large scale r-modes are also affected by this new aspect of the instability. We analyse the damping effect of shear viscosity, which should be particularly efficient at small scales, arguing that it will not be sufficient to completely suppress the instability  in astrophysical systems. ",
        "introduction": "Since the first observation of a glitch in the spindown of the Vela Pulsar in 1969 \\cite{vela1,vela2}, there have been a large number of similar events observed in other pulsars \\cite{glitch,espinoza}. Several competing,  in some cases complimentary, mechanisms have been suggested as explanation for these occurrences \\cite{sidery,pizzochero,warsaw1, warsaw2,haskell}. The two most widely classes models relate either to changes in the star's elastic crust or the dynamics of  the superfluid neutrons that are present both in the core and  the inner crust. The first set of models involve fractures in the crust, leading to changes in the moment of inertia of the star \\cite{crust}. The second set involves the unpinning of vortices associated with a superfluid component from the star's inner crust \\cite{pinned}. The unpinning event is attributed to the build-up of ``differential rotation'' between the two components (elastic crust and interpenetrating superfluid). It is the second of these two mechanisms with which we concern ourselves in this paper.  We consider a promising mechanism for triggering the observed events; a superfluid instability present in systems with ``differential rotation'', like the lag that plays a central role in all vortex-based glitch models. The basic mechanism has already been discussed in \\cite{GAprl}, where it was argued that the instability may trigger large-scale vortex unpinning leading to the observed events. The first part of the argument was a demonstration that  unstable r-modes could be generated by the difference in rotation of the charged particles (protons and electrons) and the superfluid neutrons in the star. In particular, it was shown in \\cite{GAprl} that such unstable modes exist for a strongly coupled system. The second part of the argument consisted of a demonstration that viscosity would not suppress this instability completely, but would allow growing modes to exist for a range of wavelengths once a critical rotational lag was reached. The third part of the argument placed the mechanism in an astrophysical context by comparing the predictions of the model to observational data.  In this paper we consider this new r-mode instability in more detail. We extend the analysis of \\cite{GAprl} beyond the strong-drag limit, and show that  dynamically unstable modes (with a growth time similar to the star's rotation rate) exist for a  wide range of parameter values. In addition, we discover that the r-modes also suffer a secular instability (with growth time proportional to the mutual friction) throughout much of parameter space. This aspect is new. It is particularly interesting as it affects also the  global scale r-modes, while the dynamical instability is restricted to small scale modes. As in the previous work, we limit ourselves to discussion of the non-relativistic case. Again following the approach of \\cite{GAprl}, we introduce shear viscosity and demonstrate that the instability is not completely suppressed by the inclusion of the associated damping.  The main purpose of this investigation is to establish the robustness of the superfluid instability, highlight the key ingredients that lead to its presence and set the stage for more detailed studies of this mechanism. The instability that we consider may be generic (belonging to the class of two-stream instabilities that are known to operate in a wide variety of physical systems \\cite{twostream}), but at this point it is not clear how it is affected by  other aspects of neutron star physics. In particular, future work needs to extend the analysis to account for the crust elasticity and the presence of the star's magnetic field. ",
        "conclusions": ""
    },
    "1208/1208.2845_arXiv.txt": {
        "abstract": "{ The Virgo cluster of galaxies provides excellent conditions for studying interactions of galaxies with the cluster environment. Both the high-velocity tidal interactions and effects of ram pressure stripping by the intracluster gas can be investigated in detail. } { We extend our systematic search for possible anomalies in the magnetic field structures of Virgo cluster spirals in order to characterize a variety of effects and attribute them to different disturbing agents. } { Six angularly large Virgo cluster spiral galaxies (NGC\\,4192, NGC\\,4302, NGC\\,4303, NGC\\,4321, NGC\\,4388, and NGC\\,4535) were targets of a sensitive total power and polarization study using the 100-m radio telescope in Effelsberg at 4.85~GHz and 8.35~GHz (except for NGC\\,4388 observed only at 4.85~GHz, and NGC\\,4535 observed only at 8.35~GHz). The presented two-frequency studies allow Faraday rotation analysis. } { Magnetic field structures distorted to various extent are found in all galaxies. Three galaxies (NGC\\,4302, NGC\\,4303, and NGC\\,4321) show some signs of possible tidal interactions, while NGC\\,4388 and NGC\\,4535 have very likely experienced strong ram-pressure and shearing effects, respectively, visible as distortions and asymmetries of polarized intensity distributions. As in our previous study, even strongly perturbed galaxies closely follow the radio-far-infrared correlation. In NGC\\,4303 and NGC\\,4321, we observe symmetric spiral patterns of the magnetic field and in NGC\\,4535 an asymmetric pattern. } { The cluster environment clearly affects the evolution of its member galaxies via various effects. Magnetic fields allow us to trace even weak interactions that are difficult to detect with other observations. Our results show that the degree of distortions of a galaxy is not a simple function of the distance to the cluster center but reflects also the history of its interactions. The angle $\\Theta$ between the velocity vector $\\vec{v}$ and the rotation vector $\\vec{\\Omega}$ of a galaxy may be a general parameter that describes the level of distortions of galactic magnetic fields. Information about the motions of galaxies in the sky plane and their three-dimensional distribution, as well as information about the intracluster medium can also be obtained from the Faraday rotation analysis. } ",
        "introduction": "\\label{intro}  In our previous study (We\\.zgowiec et al.~\\cite{wezgowiec}) of radio polarization in Virgo cluster spiral galaxies, we found a variety of environmental effects manifested as distorted structures of the galactic magnetic fields. These interactions have been previously investigated in \\ion{H}{i} (Cayatte et al.~\\cite{cayatte}, \\cite{cayatte2}, \\cite{viva}) and the ionised medium (Chemin et al.~\\cite{chemin}). Both data sets suggest that the environment significantly affects the interstellar medium (ISM) of cluster galaxies. Theoretical estimations and simulations of ram-ressure stripping effects have been performed by Vollmer et al.~(\\cite{vollmer2}), Schulz \\& Struck~(\\cite{schustr}), and Roediger \\& Br\\\"uggen~(\\cite{robru}). Our observations confirmed that cluster galaxies interact with their environment, as well as proved that radio polarimetry is a very sensitive tool for studying these interactions. Only few galaxies have been previously observed (see Chy\\.zy et al.~\\cite{chyzy}, Soida et al.~\\cite{soida}, Vollmer et al.~\\cite{vollmer04b}, Chy\\.zy et al.~\\cite{chyzy06}) and our present survey extends the database of magnetic properties of Virgo cluster spirals.  To enrich the quality of the statistical analysis and thus to allow studies of global trends, we extended our preliminary results presented in We\\.zgowiec et al.~(\\cite{wezgowiec}) to obtain a statistically complete sample of Virgo cluster spiral galaxies. We therefore searched for galaxies with an optical diameter of at least 4.5 arcmin (at the assumed distance to the Virgo cluster of 17\\,Mpc) and a radio flux density of the emission from the disk exceeding 40 mJy at 1.4~GHz. This selection yielded 12 galaxies in total. Five of them (NGC\\,4438, NGC\\,4501, NGC\\,4535, NGC\\,4548, and NGC\\,4654) had been observed previously (We\\.zgowiec et al.~\\cite{wezgowiec}) and two more, NGC\\,4254 and NGC\\,4569, had been observed in other projects (Chy\\.zy et al.~\\cite{chyzy07} and \\cite{chyzy06}, respectively). The remaining five galaxies (NGC\\,4192, NGC\\,4302, NGC\\,4303, NGC\\,4321, and NGC\\,4388) are the subject of this study. We also complete the observations of one galaxy from our previous sample -- NGC\\,4535 -- at 8.35~GHz providing data of higher resolution to study in greater detail the unusual asymmetry of polarized intensity found by us with sensitive low-resolution 4.85~GHz observations (see We\\.zgowiec et al.~\\cite{wezgowiec}).  Our observations of the second part of the complete sample were designed to search for cases of galaxy-cluster interactions in different parts of the Virgo cluster. This would help us to determine the extent to which the magnetic field distortions observed in our first study represent a general rule. Our present study will be used to construct a comprehensive database of magnetic structures of cluster spiral galaxies. Together with the data in other wavelengths it also will provide important clues about both the evolution and interactions of galaxies with the cluster environment. We also carry out X-ray studies of the Virgo cluster galaxies, the results of which will place tight constraints on the nature of the interactions of galaxies with the intracluster medium (ICM). Our first results were published in We\\.zgowiec et al.~(\\cite{machcones}), where we report the first detection of a hot gas halo around a Virgo cluster galaxy reminiscent of a Mach cone geometry. The thorough analysis of X-ray data from selected Virgo cluster galaxies will be presented in separate papers. ",
        "conclusions": "We have presented results of the second part of our systematic study of the magnetic-field structures of Virgo cluster spirals (see also We\\.zgowiec et al.~\\cite{wezgowiec}). As previously, we used the Effelsberg radio telescope at 4.85~GHz to detect the weak, extended total power and polarized emission, and particular objects were studied in more detail at 8.35~GHz. Our studies yielded the following results:  \\begin{itemize} \\item[-] In the edge-on galaxies NGC\\,4192 and NGC\\,4388, we found asymmetric halos, which in the latter galaxy is heavily distorted. \\item[-] NGC\\,4302 likely forms a physical pair with a close companion NGC\\,4298. The asymmetries of the polarized intensity together with H$\\alpha$ asymmetries and $\\ion{H}{i}$ tails support this scenario. Between the galaxies, there is a region of polarized emission that, as in the case of NGC\\,4038/4039, indicates that the magnetic field has been amplified by tidal interactions between both objects. \\item[-] Our higher resolution observations of NGC\\,4535 presented here confirm the strong asymmetry of the polarized emission and enhancement of the magnetic field found in We\\.zgowiec et al.~(\\cite{wezgowiec}). Since this feature lies within the extended $\\ion{H}{i}$ envelope (see Vollmer et al.~\\cite{4535shear}, it is most likely caused by shearing forces. \\item[-] Two outskirts face-on galaxies have different radio-emission distributions. While NGC\\,4303 seems completely unperturbed, NGC\\,4321 shows signs of perturbations. This may mean that the degree of distortion of a galaxy is not a simple function of its distance from the cluster center. \\item[-] The magnetic fields of both NGC\\,4303 and NGC\\,4321 have symmetric spiral patterns, whereas that of NGC\\,4535 is highly asymmetric. \\item[-] The perturbations of the galactic sttelar or ISM component or both do not significantly affect the radio -- far infrared correlation. The only exception in our sample introduced in this paper was NGC\\,4388, which displays only a slight excess of radio emission most likely due to active galactic nucleus (AGN) activity. In We\\.zgowiec et al.(~\\cite{wezgowiec}), we found a more significant excess in another AGN host -- NGC\\,4438. \\item[-] The angle $\\Theta$ between the velocity vector $\\vec{v}$ and the rotation vector $\\vec{\\Omega}$ of a galaxy may be a general parameter that describes the level of distortions in the galactic magnetic fields. \\end{itemize}  The ''magnetic diagnostics'' presented in this paper have been proven to trace environmental effects experienced by cluster galaxies interacting with either each other or the hot ICM. It provides us with a very sensitive tool for examining perturbations sometimes not yet visible in other domains. Even low-resolution but sensitive observations of the polarized intensity may yield some clues about the magnetic-field configuration via a rotation-measure analysis. Encouraged by the results of this and our previous work, we plan to extend our sample even further to obtain polarized intensity data for radio-weaker galaxies, located in different parts of the cluster. This would allow us to perform an extensive statistical study of the effects of various kinds of interactions upon the properties of galaxies. This will also be done by combining our data with available data in X-rays, $\\ion{H}{i}$, H$\\alpha$, and CO. We expect this study to provide comprehensive insight into the evolution of galaxies in a particularly ''influential'' environment."
    },
    "1208/1208.5712_arXiv.txt": {
        "abstract": "{} {Stellar mass distribution in the Andromeda galaxy (M\\,31) is estimated using optical and near-infrared imaging data. Combining the derived stellar mass model with various kinematical data, properties of the dark matter (DM) halo of the galaxy are constrained.} {SDSS observations through the $ugriz$ filters and the Spitzer imaging at 3.6 microns are used to sample the spectral energy distribution (SED) of the galaxy at each imaging pixel. Intrinsic dust extinction effects are taken into account by using far-infrared observations. Synthetic SEDs created with different stellar population synthesis models are fitted to the observed SEDs, providing estimates for the stellar mass surface density at each pixel. The stellar mass distribution of the galaxy is described with a 3-dimensional model consisting of a nucleus, a bulge, a disc, a young disc and a halo component, each following the Einasto density distribution (relations between different functional forms of the Einasto density distribution are given in Appendix~\\ref{app:2}). By comparing the stellar mass distribution to the observed rotation curve and kinematics of outer globular clusters and satellite galaxies, the DM halo parameters are estimated.} {Stellar population synthesis models suggest that M\\,31 is dominated by old ($\\ga7$\\,Gyr) stars throughout the galaxy, with the lower limit for the stellar mass-to-light ratios $M/L_r\\ga 4\\,M_{\\odot}/L_{\\odot}$. The upper limit $M/L_r\\la 6\\,M_{\\odot}/L_{\\odot}$ is given by the rotation curve of the galaxy. The total stellar mass is (10--$15)\\cdot10^{10}M_{\\odot}$, 30\\% of which is in the bulge and 56\\% in the disc. None of the tested DM distribution models (Einasto, NFW, Moore, Burkert) can be falsified on the basis of the stellar matter distribution and the rotation curve of the galaxy. The virial mass $M_{200}$ of the DM halo is $(0.8$--$1.1)\\cdot10^{12}M_{\\odot}$ and the virial radius is $R_{200}=$189--213~kpc, depending on the DM distribution. For the Einasto profile, the average density of the DM halo within the central 10 pc is 16--61\\,$M_{\\odot}\\mathrm{\\,pc^{-3}}$ (0.6--2.3\\,$\\mathrm{TeV}/c^2\\,\\mathrm{cm}^{-3})$, depending on the stellar mass model. The central density of the DM halo is comparable to that of nearby dwarf galaxies, low-surface-brightness galaxies and distant massive disc galaxies, thus the evolution of central DM halo properties seems to be regulated by similar processes for a broad range of halo masses, environments, and cosmological epochs.} {} ",
        "introduction": "Due to its proximity and size, our nearest large neighbour galaxy M\\,31 offers a unique and attractive opportunity to study stellar populations and galactic structure in detail. It has been a source for groundbreaking discoveries ever since the secure acknowledgement of this nebula as an extragalactic object by Ernst \\\"{O}pik \\citeyearpar{opik:22}.  Self-consistent treatment of the available photometry and kinematical data of M\\,31 enabled the construction of sophisticated multi-component galactic models already decades ago \\citep[e.g.][]{tenjes:94}. Considering the huge volume and high detail of observational information available nowadays, complex mass models of M\\,31 offer a promising opportunity for casting light on one of the most puzzling problems in astrophysics and cosmology: the nature and properties of dark matter (DM) haloes. By now, the increasing scope of observations has enabled stretching mass models much further than the extent of gas disc rotation curves, providing new clues about DM halo parameters \\citep{geehan:06,seigar:08,chemin:09,corbelli:10}.  On the other hand, particle physics instrumentation has reached a level at which it can provide some hints about DM. Although the diffuse Galactic background likely exceeds the expected flux of high-energy particles resulting from decaying or annihilating DM in extragalactic sources \\citep{bertone:07,Hutsi:10}, particles from more concentrated extragalactic objects might be detectable as an enhancement of the Galactic signal within certain apertures. By comparing the assumed DM distribution in M\\,31 to the data collected with the diverse arsenal of ground-based and space-borne detectors of high-energy particles, some constraints on the energy spectrum of DM particles have already been laid \\citep[e.g.][]{aharonian:03,lavalle:06,Boyarsky:08,dugger:10,Watson:12}.  For more stringent constraints, not only more capable detectors are needed but also a better understanding of the properties of the source DM haloes. As bizarre as it seems, the derivation of the detailed mass distribution of the \\object{Andromeda} galaxy was limited by the lack of suitable optical imaging up to recent times.  Although visible even to the naked eye, its span over four degrees on the celestial sphere makes Andromeda a real challenge to observe with a usual scientific telescope. Thus it is only very recently that observations covering the galaxy with deep wide-field CCD imaging have started to appear: a dedicated scan within the Sloan Digital Sky Survey \\citep{york:00}, the Canada-France-Hawaii telescope Megacam programme PAndAS \\citep{mcconnachie:09}, the Pan-STARRS telescope project PAndromeda \\citep[][]{lee:12}. Combined with the space-based ultraviolet \\citep{thilker:05}, near- \\citep{barmby:06} and far-infrared \\citep{gordon:06,fritz:11} observations, these data now provide an unprecedented panchromatic view of a galaxy at a resolution of a few parsecs, allowing for the derivation of detailed properties of the stellar populations.  In this paper we estimate the stellar and DM distribution of M\\,31 using the SDSS and Spitzer 3.6-micron imaging to constrain the properties of the stellar populations. We construct a mass distribution model of the galaxy in correspondence with the latest kinematical data from the literature, giving estimates for the DM halo properties and the related uncertainties.  We have taken the inclination angle of M\\,31 to be 77.5\\degr \\citep{Walterbos:88,deVaucouleurs:91} and the distance to the galaxy \\mbox{785\\,kpc} \\citep{McConnachie:05}, yielding the scale \\mbox{228\\,pc/arcmin}. Throughout the paper, luminosities are presented in AB-magnitudes and are corrected for the Galactic extinction according to \\citet{tempel:11}, where extinction corresponding to the Sloan filters is derived from the \\citet{Schlegel:98} estimates and the Galactic extinction law by linear interpolation. The absolute solar luminosity for each filter is taken from \\citet{blanton:07}. ",
        "conclusions": "\\begin{table*} \\caption{Comparison of bulge, disc, and DM halo mass estimates. All masses are in $10^{10}M_\\odot$.} \\label{table:mass_comp} \\centering \\begin{tabular}{lccc} \\hline\\hline Model &  $M_{\\mathrm{bulge}}$ &  $M_{\\mathrm{disc}}$   &  $M_{200}$  \\\\ \\hline \\citet{geehan:06}, best-fit (maximum-disc) model    & 3.3       & 8.4 (13.7) &  68 (94)                 \\\\ \\citet{seigar:08}, model without (with) adiabatic contraction   & 3.5 (3.5) & 5.8 (7.3)  & 73 (89)  \\\\ \\citet{chemin:09}, ``hybrid\" model    & 2.32      & 7.1        &  100                     \\\\ \\citet{corbelli:10}, NFW model with $(M/L)_\\mathrm{bulge} = (M/L)_\\mathrm{disc}$  & 3.8 & 8.8 &  85\\tablefootmark{c} \\\\ This work, B07 model                            & 4.4\\tablefootmark{a}   &  5.7\\tablefootmark{b}   &  113              \\\\ This work, maximum-stellar model\t\t      & 6.6\\tablefootmark{a} & 8.6\\tablefootmark{b} &  127\\\\ \\hline \\end{tabular} \\tablefoot{ \\tablefoottext{a}{Sum of the bulge and stellar halo masses.} \\tablefoottext{b}{Sum of the disc and young disc masses.} \\tablefoottext{c}{Recalculated from $M_{98}$ in the original paper.} } \\end{table*}  \\begin{figure} \\includegraphics[width=88mm]{figur10.eps} \\caption{Average DM density inside a given radius, corresponding to different DM distributions in the case of the B07 stellar masses. For the Einasto DM distribution, also the maximum-stellar-mass case is plotted. For comparison, central densities of some nearby dwarf galaxies and low-surface-brightness galaxies are shown. The triangular/quadrangular datapoints are calculated assuming the Burkert DM, the circular datapoints correspond to the NFW DM.} \\label{fig:dm_cumul} \\end{figure}   Let us compare our mass models of M\\,31 to some other recently constructed models. In Table~\\ref{table:mass_comp}, mass estimates suggested by our models for the bulge, disc, and DM halo are compared to the estimates by \\citet{geehan:06}, \\citet{seigar:08}, \\citet{chemin:09}, and \\citet{corbelli:10}. For a better understanding of the compatibility of these models, we will briefly summarise the principal properties and differences of these models below.  The model derived by \\citet{geehan:06} consists of a central supermassive black hole, a bulge, a disc, and a DM halo. Its stellar components are determined using various luminosity measurements in $V$, $R$, and $r$ filters out to the (projected) distance of 25\\,kpc along the major axis, and the kinematics is based on a composite rotation curve. Mass-to-light ratios of the stellar components are treated as free parameters. The total mass is constrained by data on the motions of outer planetary nebulae, globular clusters and satellite galaxies. In Table~\\ref{table:mass_comp}, \\citet{geehan:06} values for the best-fit model and the maximum-disc model (in brackets) are presented.  \\citet{seigar:08} constructed a similar black hole\\,+\\,bulge\\,+\\,disc\\,+\\,DM halo model. They used the Spitzer 3.6-micron imaging data and \\mbox{$B\\!-\\!R$} colour profile to determine the mass-to-light ratios; dynamical mass estimators were the same or similar as in \\citet{geehan:06}. In addition to the usual DM profiles, \\citet{seigar:08} considered the case of a dark halo that has undergone an adiabatic contraction due to the gravitational attraction of the baryonic material. In Table~\\ref{table:mass_comp}, the model with adiabatic contraction is presented in brackets; it should not be compared directly to the other models.  \\citet{chemin:09} constructed a black hole\\,+\\,bulge\\,+\\,disc\\,+\\,gas model, using their newer \\ion{H}{i} data for constraining the kinematics. For a more accurate description of the disc potential, the disc density distribution was settled as the residual of the surface brightness distribution after the subtraction of the bulge contribution. The small contribution of gas mass was considered on the basis of $\\mathrm{H}_2$ and \\ion{H}{i} surveys. In Table~\\ref{table:mass_comp}, the ``hybrid\" model (with the bulge mass is determined from stellar velocity dispersions and the disc mass from simple stellar population models) values of \\citet{chemin:09} are presented.  \\citet{corbelli:10} used a bulge\\,+\\,disc\\,+\\,gas model together with an optional Burkert/NFW DM halo to fit their \\ion{H}{i} kinematics data and outer dynamics estimators. The best-fit model with equal bulge and disc mass-to-light ratios and the NFW dark halo is shown in Table~\\ref{table:mass_comp}. We have rescaled the virial mass $M_{98}$ given by \\citet{corbelli:10} to $M_{200}$.  In contrast to these four works, our stellar model is based on fully 2-dimensional dust-corrected imaging through 6 filters and some more recent dynamical mass estimators. Instead of the central black hole, our model includes the nucleus of the galaxy, which contributes significantly more to the total mass and the model rotation curve of the galaxy. This contribution is nevertheless tiny and has a negligible effect on other model parameters, as well as our usage of an oblate bulge \\citep[with an axial ratio 0.8;][]{tempel:11} instead of a spherical one. In Table~\\ref{table:mass_comp}, our B07 model and the maximum-stellar model results are given. The actual values probably lie between the estimates of these two models.  Table~\\ref{table:mass_comp} shows that the bulge mass suggested by our models is somewhat higher than in previous models, probably resulting from different bulge parameters but also because we have used a larger set of observational data to constrain the stellar populations. Nevertheless, the bulge dominates the total gravitational potential only up to the radius 6--8\\,kpc and bulge properties have little effect on the DM halo parameters.  As can be seen from Table~\\ref{table:mass_comp}, the DM halo virial mass estimates have previously remained mostly below 1$\\cdot10^{12}M_\\odot$, whereas our models suggest slightly higher values, (1.0--1.3)\\,$\\cdot10^{12}M_\\odot$. Once again, the most likely source of differences is our usage of a larger collection of observational data on the outer dynamics. As shown in Figs.~\\ref{fig:dm_params} and~\\ref{fig:app_dm_params}, the virial mass is almost independent of the DM density profile and the stellar mass model, being uniquely determined by the outer dynamics of the galaxy.  In Fig.~\\ref{fig:dm_cumul}, the distribution of the average DM halo density within a given radius is shown for each DM density distribution model. The added datapoints provide an illustrative comparison with DM haloes of other galaxies, for which the average central density has been recently measured more or less reliably: nearby dwarf galaxies \\citep{Gilmore:07, Gentile:07a, Gentile:07b, Oh:11, Adams:12, Breddels:12} and low-surface-brightness galaxies \\citep{Coccato:08, deBlok:08, KuziodeNaray:08}. Some of these values are calculated for the modified-isothermal DM distribution, others for the NFW distribution. In several cases, both versions are presented. These datapoints should be compared to the Burkert and NFW profiles of M\\,31, respectively. It is seen that the central density of DM haloes varies by a couple of orders of magnitude and despite its higher total mass, the DM halo of M\\,31 cannot be distinguished from an average dwarf or low-surface-brightness galaxy in this aspect. Interestingly, the estimate of the central density (0.012--0.028)$\\cdot M_{\\odot}\\mathrm{pc^{-3}}$ of DM haloes of massive disc galaxies near redshift $z\\simeq0.9$ \\citep{tamm:05} also falls within this range, hinting that the DM halo concentration process seems to be restricted rather uniformly over a very wide variety of halo masses, environments, and cosmological epochs.  To conclude our work, it is interesting and also disappointing to note that the usage of additional observational data does not reduce significantly the uncertainties and scatter of the parameters of M\\,31 mass distribution models. Our vague understanding of the evolution of the physical properties of stellar populations restrains the gain from all the gathered observational information on the chemical content and formation history of a galaxy. Despite the improved imaging and kinematics data, we are still unable to confirm or rule out the maximum-disc or maximum-baryonic approach in splitting the contributions of luminous and dark matter to the overall mass distribution. It can only be concluded that the bulge mass of M\\,31 probably lies within the range (4.4--6.6)\\,$\\cdot10^{10}M_\\odot$ and the disc mass within the range (5.7--8.6)\\,$\\cdot10^{10}M_\\odot$. Nevertheless, M\\,31 provides an exceptional opportunity to estimate the virial mass and the outer distribution of a DM halo thanks to the possibility of tracing the gravitational potential with various test bodies at large distances from the galactic centre."
    },
    "1208/1208.0906_arXiv.txt": {
        "abstract": "Two multifrequency campaigns were carried out on OJ287 in 2005: in April when it was in its pre-outburst state, and in November, during the main 12 yr cycle outburst. The wavelength coverage was from radio to X-rays. In the optical-to-UV range the differential spectrum between the observations has a bremsstrahlung spectral shape, consistent with gas at $3 \\times 10^{5}K$ temperature. Our result supports the hydrogen column density of the OJ287 host galaxy of $\\sim9.3\\times 10^{20} cm^{-2}$, the average value found by Gosh \\& Soundararajaperumal. The $3 \\times 10^{5}K$ bremsstrahlung radiation was predicted in the binary black hole model of OJ287, and it arises from a hot bubble of gas which is torn off the accretion disc by the impact of the secondary. As this radiation is not Doppler boosted, the brightness of the outburst provides an estimate for the mass of the secondary black hole, $\\sim1.4\\times10^{8}$ solar mass. In order to estimate the mass of the primary black hole, we ask what is the minimum mass ratio in a binary system which allows the stability of the accretion disc. By using particle simulations, we find that the ratio is $\\sim1.3\\times10^{2}$. This makes the minimum mass of the primary $\\sim1.8\\times10^{10}$ solar mass, in agreement with the mass determined from the orbit solution, $1.84 \\times 10^{10}$ solar mass. With this mass value and the measured K-magnitude of the bulge of the host galaxy of OJ287, the system lies almost exactly on the previously established correlation in the black hole mass vs. K-magnitude diagramme. It supports the extension of this correlation to brighter magnitudes and to more massive black holes than has been done previously. ",
        "introduction": "There is a fundamental relation between the central black hole mass and bulge luminosity in galaxies (Magorrian et al. 1998, Ferrarese \\& Merritt 2000, G\\\"ultekin et al. 2009, Kormendy and Bender 2011a,b). In the study of the tightness of this correlation one should use the most accurate mass values based on the dynamics of bodies orbiting the central black hole. So far there are not many such mass determinations. The mass of the black hole in the Galactic centre gives a fixed point at the low end of the scale (Ghez et al. 2003, Genzel et al. 2010), the masses of megamaser sources fall in the middle of the range (e.g. NGC 4258, Miyoshi et al. 1995, Kuo et al. 2011), while the mass of the OJ287 primary gives the only point at the high mass end (Valtonen et al. 2010a). In addition, there are plenty of data based on the velocity dispersions of galactic centres (e.g M87, Gebhardt \\& Thomas 2009), reverberation mapping measurements and finally much recent data based on emission line widths in quasars (for a review, see Ferrarese \\& Ford 2005). These may be viewed as giving reliable but less accurate mass values (Valluri et al. 2004, Denney et al. 2009, 2010, Kelly et al. 2010). Even though the single data point at the high mass end is very accurate, with one percent uncertainty, one would like to have independent evidence that it is at least approximately correct. Then one could use the dynamical mass with greater confidence in correlation diagrams. Thus the primary objective in this paper is to confirm the value of the central black hole mass in OJ287.  OJ287 has a historical optical light curve extending over 120 yr from late 1800's up to date. It shows a prominent twelve year cycle (Sillanp\\\"a\\\"a et al. 1988, Valtonen et al. 2006a), with an additional two peak structure with 1 - 2 yr peak separation (Valtonen et al. 2010b). It was suggested by Sillanp\\\"a\\\"a et al. (1988) that the 12 yr structure arises from a black hole binary system where a secondary perturbs the accretion disc of the primary at regular intervals and causes indirectly increased emission in the jet of OJ287. The jet is lined up with the observer at a small viewing angle which further enforces the amplitude of brightness variations via Doppler boosting (L\\\"ahteenm\\\"aki \\& Valtaoja 1999, Jorstad et al. 2005). It is extremely unlikely ($\\chi^2$ probability less than $10^{-8}$) that that a single black hole model would produce such good light curve predictions as the binary model has done in the past 14 years (Valtonen et al. 2011).  This model lead to the prediction that the next major outburst after 1983 should take place in the autumn of 1994. The outburst came as scheduled (Sillanp\\\"a\\\"a et al. 1996). At this point several more detailed binary black hole models were put forward, with differing predictions for the next major outburst (see Kidger 2000). These were in two general categories: exactly periodic (non-precessing) models which predicted the main outburst in the autumn of 2006, and precessing models which gave the outburst time a year earlier (Lehto \\& Valtonen 1996, Sundelius et al. 1997). There are two reasons for this. First, in the latter models the outbursts are related to the impacts of the secondary on the accretion disc of the primary. With forward precession, as in General Relativity, the impact times happen earlier than in a non-precessing periodic model. This makes about $\\sim 6$ month contribution to the advance of the outburst time (Lehto \\& Valtonen 1996). The second contribution comes from the fact that the accretion disc is lifted up by the approaching secondary (Sundelius et al. 1996, Ivanov et al. 1998) and the impact happens earlier than expected in a rigid disc model by another $\\sim 6$ months.  In order to vertify the correctness, or otherwise, of these two clearly different models, a major monitoring campaign was set up, with frequent observations starting in early 2005. Not only was the timing important, but also the nature of the radiation at the outburst peak. Most periodic models were predicting an increase of radiation from the jet, i.e. the outburst should have a power-law like synchrotron spectrum. In contrast, the precessing binary model predicted that the radiation is bremsstrahlung from hot gas ($1 - 4\\times10^{5} K$, the exact value is calculated below) which has been torn off the accretion disc. Thus, it was important to study the spectral energy distribution during the outburst, and to compare it with the spectral distribution outside the outburst time. We had two periods of XMM-Newton observations, one in April 2005 prior to the outburst, and another one in November 2005 when the outburst was expected in the precessing model. Preliminary reports of these observations have been given in Ciprini et al. (2007), Ciprini and Rizzi (2009) and Ciprini (2010). \\begin{figure} \\includegraphics[width=6cm,angle=270]{ValtonenFigure1.eps} \\caption{The optical light curve of OJ287 during the 2005 outburst. Data points are based on Valtonen et al. (2008, 2009).} \\end{figure}  As to the timing, the outburst followed the prediction of the precessing model (Valtonen et al. 2006b, Valtonen et al. 2008). The primary outburst came in 2005 October/November while OJ287 was in an exceptionally quiet state in the autumn of 2006 when the periodic peak brightness should have happened. Thus the result was unusually clear. However, it is also crucial to verify the second hypotheses of the precessing binary model, i.e. to see if the spectrum agrees with bremsstrahlung radiation. A major problem here would be the contribution from the ``normal'' background activity which produces secondary peaks of radiation. Luckily the second XMM-Newton period fell between two secondary peaks, and thus our measurements should tell about the nature of the primary outburst (Figure 1).  In this paper we first discuss the differential energy distribution between the two XMM-Newton observing periods, and confirm that the excess has likely bremsstrahlung nature. Since the radiation from the disc impact is not Doppler boosted, the strength of the outburst is directly linked to the mass of the secondary. We then ask what is the minimum mass of the primary which provides the stability for the accretion disc. This gives us the minimum value for the primary mass. Finally we compare this with the dynamical mass determination, and after being satisfied that the two methods give consistent results, produce a slightly modified black hole mass - K-magnitude correlation diagramme with OJ287 added to the older data. ",
        "conclusions": "We find that the mass values for the two black holes in the OJ287 system, as derived from the orbital solution, are quite consistent with values obtained from astrophysical arguments. The chain of arguments starts from identifying the 2005 outburst radiation as bremsstrahlung. Since its likely origin is the collision of the secondary on the accretion disc of the primary, this allows the determination of the mass of the secondary. The requirement that the disc is stable inspite of the binary action puts a lower limit on the mass of the primary. It is very near the mass determined from the orbit solution. Thus we can have increased confidence in the correctness of the orbit solution.  We place OJ287 in the  black hole mass vs. K-magnitude  diagramme and find that it lies almost exactly on the mean line determined from the lower black hole masses. This gives us confidence that this relation can be extended to up to the largest measured black hole masses.  As a side result we find that the host galaxy of OJ287 has a moderate amount of extinction which corresponds to the hydrogen column density of $\\sim9.3\\times10^{20} cm^{-2}$, similar to values measured in some other galaxies from the 2175 A extinction feature (Wang et al. 2004). The dust and the perturbed appearance of the host both suggest that OJ287 has acquired its second black hole in a merger of two galaxies, possibly leading to binary black hole evolution as calculated by Iwasawa et al. (2010).  The observed periodicities in the light curves of OJ287 are consistent with a model where both the slowly spinning primary black hole and the fast spinning secondary black hole contribute to the optical light variations (Neronov and Vovk 2011). In future it may become possible to see the orbital motion directly when the required microarcsecond resolution is achieved. It may also be possible to detect the motion of the secondary from its spectral lines. The orbital speed of the secondary is so fast that its spectral lines are shifted very far from their normal positions, outside the spectral region where the line searches have been carried out (Nilsson et al. 2010)."
    },
    "1208/1208.3522_arXiv.txt": {
        "abstract": "We study the volume-limited and nearly mass selected (stellar mass $M_{\\rm stars}\\ga6\\times10^9$ M$_\\odot$) \\atl\\ sample of 260 early-type galaxies (ETGs, ellipticals Es and lenticulars S0s). We construct detailed axisymmetric dynamical models (JAM), which allow for orbital anisotropy, include a dark matter halo, and reproduce in detail both the galaxy images and the high-quality integral-field stellar kinematics out to about 1\\re, the projected half-light radius. We derive accurate total mass-to-light ratios $(M/L)_e$ and dark matter fractions $f_{\\rm DM}$, within a sphere of radius $r=\\re$ centred on the galaxies. We also measure the stellar $(M/L)_{\\rm stars}$ and derive a median dark matter fraction $f_{\\rm DM}=13\\%$ in our sample. We infer masses $M_{\\rm JAM}\\equiv L\\times (M/L)_e \\approx 2\\times M_{1/2}$, where $M_{1/2}$ is the total mass within a sphere enclosing half of the galaxy light.  We find that the thin two-dimensional subset spanned by galaxies in the $(M_{\\rm JAM},\\sigma_e,R_e^{\\rm maj})$ coordinates system, which we call the Mass Plane (MP) has an observed rms scatter of 19\\%, which implies an intrinsic one of 11\\%. Here $R_e^{\\rm maj}$ is the major axis of an isophote enclosing half of the {\\em observed} galaxy light, while $\\sigma_e$ is measured within that isophote. The MP satisfies the scalar virial relation $M_{\\rm JAM}\\propto\\sigma_e^2 R_e^{\\rm maj}$ within our tight errors. This show that the larger scatter in the Fundamental Plane (FP) $(L,\\sigma_e,R_e)$ is due to stellar population effects (including trends in the stellar Initial Mass Function [IMF]). It confirms that the FP deviation from the virial exponents is due to a genuine $(M/L)_e$ variation. However, the details of how both \\re\\ and $\\sigma_e$ are determined are critical in defining the precise deviation from the virial exponents. The main uncertainty in masses or $M/L$ estimates using the scalar virial relation is in the measurement of \\re. This problem is already relevant for nearby galaxies and may cause significant biases in virial mass and size determinations at high-redshift. Dynamical models can eliminate these problems. We revisit the $(M/L)_e-\\sigma_e$ relation, which describes most of the deviations between the MP and the FP. The best-fitting relation is $(M/L)_e\\propto \\sigma_e^{0.72}$ ($r$-band). It provides an upper limit to any systematic increase of the IMF mass normalization with $\\sigma_e$. The correlation is more shallow and has smaller scatter for slow rotating systems or for galaxies in Virgo. For the latter, when using the best distance estimates, we observe a scatter in $(M/L)_e$ of 11\\%, and infer an intrinsic one of 8\\%. We perform an accurate empirical study of the link between $\\sigma_e$ and the galaxies circular velocity $V_{\\rm circ}$ within 1\\re\\ (where stars dominate) and find the relation $\\max(V_{\\rm circ}) \\approx 1.76\\times \\sigma_e$, which has an observed scatter of 7\\%. The accurate parameters described in this paper are used in the companion Paper~XX of this series to explore the variation of global galaxy properties, including the IMF, on the projections of the MP. ",
        "introduction": "Scaling relations of early-type galaxies (ETGs, ellipticals E and lenticulars S0) have played a central role in our understanding of galaxy evolution, since the discovery that the stellar velocity dispersion $\\sigma$ \\citep{Minkowski1962,Faber1976} and the galaxy projected half-light radius $\\re$ \\citep{Kormendy1977} correlate with galaxy luminosity $L$. An important step forward was made with the discovery that these two relations are just projections of a relatively narrow plane, the Fundamental Plane (FP) \\citep{Faber1987,Dressler1987,Djorgovski1987}, relating the three variables $(L,\\sigma_e,\\re)$. When the plane is used as a distance indicator, as was especially the case at the time of its discovery, the luminosity can be replaced by the surface brightness within \\re\\ as $\\Sigma_e\\equiv L/(2\\pi R_{\\rm e}^2)$ and the observed plane assumes the form \\begin{equation} \\re \\propto \\sigma^{1.33} \\Sigma_e^{-0.82}, \\end{equation} where the adopted parameters are the median of the 11 independent determinations tabulated in \\citet{Bernardi2003fp}.  It was immediately realized that the existence of the FP could be due to the galaxies being in virial equilibrium \\citep[e.g.][]{Binney2008} and that the deviation (tilt) of the coefficients from the virial predictions $\\re \\propto \\sigma^2 \\Sigma_e^{-1}$, could be explained by a smooth power-law variation of mass-to-light ratio $M/L$ with mass \\citep{Faber1987}. The FP showed that galaxies assemble via regular processes and that their properties are closely related to their mass. The tightness of the plane gives constraints on the variation of stellar population among galaxies of similar characteristics and on their dark matter content \\citep{Renzini1993,Borriello2003fp}. The regularity also allows one to use the FP to study galaxy evolution, by tracing its variations with redshift \\citep{vanDokkum1996}.  However, other reasons for the deviation of the coefficients are possible: the constant coefficients in the simple virial relation only rigorously apply if galaxies are spherical and homologous systems, with similar profiles and dark matter fraction. But both galaxies concentration \\citep{Caon1993} and the amount of random motions in their stars \\citep{Davies1983} were found to systematically increase with galaxy luminosity.  The uncertain origin of the tilt led to a large number of investigations about its origin, exploring the effects of (i) the systematic variation in the stellar population or IMF \\citep[e.g.][]{Prugniel1996fp,Forbes1998fp}, or (ii) the non-homology in the surface brightness distribution \\citep[e.g.][]{Prugniel1997fp,Graham1997fp,Bertin2002fp,Trujillo2004fp} or (iii) the kinematic \\citep[e.g.][]{Prugniel1994,Busarello1997}, or (iv) the variation in the amount of dark matter \\citep[e.g.][]{Renzini1993,Ciotti1996,Borriello2003fp}, on the FP tilt and scatter. Those works were all based on approximate galaxy spherical models, trying to test general hypotheses and not reproducing real galaxies in detail, which sometimes led to contrasting results. What became clear however was that various effects could potentially influence a major part of the FP tilt. Moreover it was found that the small scatter in the FP implies a well regulated formation for ETGs.  The next step forward came with subsequent studies, which instead of testing general trends, used small samples of objects and tried to push to the limit the accuracy of measuring galaxy central masses, while reducing biases as much as possible. Those accurate total masses could be directly compared to the simple virial ones, testing for residual trends. Similar but independent studies were performed using two completely different techniques, either stellar dynamics \\citep{Cappellari2006} or strong gravitational lensing \\citep{Bolton2007mp,Bolton2008,Auger2010}. The results from those efforts agree with each others, and showed that the tilt of the FP is almost entirely due to a genuine $M/L$ variation.  In this paper we investigate once more the origin of the FP tilt. This new study is motivated by the dramatic increase in the size and quality of our galaxy sample, with respect to any previous similar study. We have in fact state-of-the-art \\sauron\\ \\citep{Bacon2001} stellar kinematics for all the 260 early-type galaxies of the \\atl\\ sample \\citep[hereafter Paper~I]{Cappellari2011a}, which constitute a volume-limited and carefully selected sample of ETGs, down to a stellar mass of about $M_{\\rm stars}\\ga6\\times10^9$ M$_\\odot$. This fact, combined with detailed dynamical models for the entire sample, allows us to test previous claims with unprecedented accuracy. The new models also include a dark matter halo and give constraints on the dark matter content in the centres of early-type galaxies. These measurements will be used in the companion \\citet[hereafter Paper~XX]{Cappellari2012p20} to provide a novel view of galaxy scaling relations.  In what follows, in Section~2 we present the sample and data, in Section~3 we describe the methods used to extract our quantities, in Section~4 we present our results on the FP tilt, dark matter and the $(M/L)-\\sigma$ relation, and finally we summarize our paper in Section~5. ",
        "conclusions": "We construct detailed dynamical models (JAM), based on the Jeans equations and allowing for orbital anisotropy, for the volume-limited and essentially mass-selected \\atl\\ sample of early-type galaxies. The models fit in detail the two-dimensional galaxy images and reproduce in detail the integral-field stellar kinematics obtained with \\sauron\\ out to about 1\\re, the projected half-light radius. We derive accurate total mass-to-light ratios $(M/L)_e$ and dark matter fractions $f_{\\rm DM}$, within a sphere of radius $r=\\re$ centred on the galaxies. We infer masses $M_{\\rm JAM}\\equiv L\\times (M/L)_e \\approx 2\\times M_{1/2}$, where $M_{1/2}$ is the mass within a sphere enclosing half of the galaxy light. We also measure stellar $(M/L)_{\\rm stars}$.  We test the accuracy of our mass determinations by running models with and without dark matter and we find that the enclosed total $(M/L)_e$ is a robust quantity, independent of the inclusion of a dark-matter halo, with an rms accuracy of 5\\% and negligible bias. In other words, even using simple mass-follow-light models, one recovers the total enclosed $(M/L)_e$ with good accuracy and small bias. We illustrate the tecniques we use to measure radii and global kinematical quantities from our data, and to robustly fit linear relations or planes to the data, even in the presence of outliers and significant intrinsic scatter. We stress the difficulty of measuring  absolutely calibrated effective radii \\re, and we argue againt extrapolation in the profiles, for more reproducible results. Systematic offsets in \\re\\ determinations are the main limitation for the use of the scalar virial relation for mass estimates, and may affect size comparisons as a function of redshift.  We find that the thin two-dimensional subset spanned by galaxies in the $(M_{\\rm JAM},\\sigma_e,R_e^{\\rm maj})$ coordinates system, which we call the Mass Plane (MP) has an observed rms scatter of 19\\%, which would imply an intrinsic one of just 11\\%. The MP satisfies the scalar virial relation $M_{\\rm JAM}\\propto\\sigma_e^2 R_e^{\\rm maj}$ within our tight errors. However, this is only true if one pays special attention to the methodology employed to determine the galaxy global parameters and in particular, (i) one uses as scale radius the major axis $ R_{\\rm e}^{\\rm maj}$ of the `effective' isophote enclosing half of the total projected galaxy light (without extrapolating the profile beyond the data), and (ii) one measures the velocity dispersion $\\sigma_e$ (which includes rotation and random motions) from a spectrum derived inside that effective isophote. This confirms with unprecedented accuracy previous claims \\citep{Cappellari2006,Bolton2008} that galaxies accurately satisfy the virial relations and that the existence of the FP is entirely explained by virial equilibrium plus a systematic variation in the total $(M/L)_e$.  We revisit the $(M/L)_e-\\sigma$ relation and measure a marginally shallower observed slope than previously reported. The minor difference can be explained by selection of the sample of galaxies previously used to fit the relations. We find that the correlation depends both on galaxy rotation and environment, in the sense that both for the subsamples of the galaxies in Virgo, or for the subsample of slow rotators,  the relation is more shallow and has a reduced scatter. In the best case, when the most accurate distances are used, the observed scatter drops to 11\\% and the intrinsic one is estimated to be a mere 8\\%.  We study the correlation between $\\sigma_e$ and the circular velocity from the dynamical models. We find that $V_{\\rm circ}(R_e^{\\rm maj}) \\approx 1.51\\times \\sigma_e$ and $\\max(V_{\\rm circ}) \\approx 1.76\\times \\sigma_e$. The relations have an observed scatter of 7--8\\% and the coefficient is independent on $\\sigma_e$.  The accurate global dynamical scaling parameters for the ETGs in the \\atl\\ sample are used in the companion Paper~XX to explore different projection of the Mass Plane and the variation of galaxy physical parameters."
    },
    "1208/1208.1527_arXiv.txt": {
        "abstract": "We explore the spatio-temporal evolution of solar flares by fitting a radial expansion model $r(t)$ that consists of an exponentially growing acceleration phase, followed by a deceleration phase that is parameterized by the generalized diffusion function $r(t) \\propto \\kappa (t-t_1)^{\\beta/2}$, which includes the logistic growth limit ($\\beta=0$), sub-diffusion ($\\beta = 0-1$), classical diffusion ($\\beta=1$), super-diffusion ($\\beta = 1-2$), and the linear expansion limit ($\\beta=2$). We analyze all M and X-class flares observed with GOES and AIA/SDO during the first two years of the SDO mission, amounting to 155 events. We find that most flares operate in the sub-diffusive regime ($\\beta=0.53\\pm0.27$), which we interpret in terms of anisotropic chain reactions of intermittent magnetic reconnection episodes in a low plasma-$\\beta$ corona. We find a mean propagation speed of $v=15\\pm12$ km s$^{-1}$, with maximum speeds of $v_{max}=80 \\pm 85$ km s$^{-1}$ per flare, which is substantially slower than the sonic speeds expected for thermal diffusion of flare plasmas. The diffusive characteristics established here (for the first time for solar flares) is consistent with the fractal-diffusive self-organized criticality (FD-SOC) model, which predicted diffusive transport merely based on cellular automaton simulations. ",
        "introduction": "Many physical processes can be characterized by their spatio-temporal evolution, which we simply define here as the temporal evolution or time dependence of a spatial or geometric parameter, say $x(t)$, where the geometric parameter $x$ could be a spatial length scale, an area, a volume, a fractal dimension, or any combination of these. At the largest scales, for instance, the Big Bang theory describes the evolution of the universe by the gradual expansion of its radius $r(t)$, which can be decelerating in a closed universe (Friedmann-Lema\\^itre-Robertson-Walker model). Recent supernova observations reveal an accelerating universe at $z \\approx 0.5$ (Riess et al.~1998; Perlmutter et al.~1999), and a deceleration that preceded the current epoch of cosmic acceleration (Riess et al.~2004).  Other dynamic processes that predict a specific spatio-temporal evolution based on some physical model include Brownian motion, fractal Brownian motion, classical diffusion, sub-diffusion, super-diffusion, L\\'evy flights, logistic growth, percolation, self-organized criticality avalanches, cellular automatons, non-extensive Tsallis entropy, complex networks, etc. The knowledge of the spatio-temporal evolution of a physical process is often directly related to scaling laws between various physical parameters, and thus spatio-temporal measurements play a decisive role in the derivation of scaling laws. For instance, particle conservation ($n V = const$) in an expanding gas, plasma, or universe predicts a reciprocal scaling between the particle density $n(t)$ and the volume $V(t)$, i.e., $n(t) \\propto V(t)^{-1}$, and consequently a scaling of $n(t) \\propto r(t)^{-3}$ as a function of the radius $r(t)$ in the case of homogeneous isotropic expansion, while it scales as $n(t) \\propto r(t)^{-D}$ for a fractal volume with Hausdorff dimension $D < 3$. The existence of a spatio-temporal evolution, such as $n(t) \\propto r(t)^{-3}$ for adiabatic expansion, implies then also the prediction of a statistical correlation or scaling law $n \\propto r^{-3}$, if the expansion speeds $v = \\partial r / \\partial t$ of the sample have a limited range.  Needless to say, that scaling laws obtained from solar data, where we have ample spatial resolution, are extremely useful for the interpretation of stellar data, where we have no spatial resolution at all and have to rely on scaling laws measured in solar or magnetospheric plasmas. Here we focus on the spatio-temporal evolution of solar flares, which is a completely unexplored topic, but bears important information on the underlying dynamic processes. There are very few statistical measurements of spatial scales $L$, areas $A$, and volumes $V$, or fractal dimensions $D$ of solar flares, and virtually no statistical studies about the temporal evolution of these parameters, such as $L(t), A(t), V(t), D(t)$. A few statistical measurements of spatio-temporal parameters of solar flares (compiled in Aschwanden 1999) have been made from the S-054 soft X-ray imager onboard {\\sl Skylab} (Pallavicini et al.~1977), from {\\sl Yohkoh} soft X-ray images (Kano and Tsuneta 1995; Porter and Klimchuk 1995; Aschwanden et al.~1996; Metcalf and Fisher 1996; Reale et al.~1997; Shimizu 1997; Garcia 1998; Nagashima and Yokoyama 2006), from the {\\sl Multispectral Solar Telescope Array (MTSA)} rocket flight (Kankelborg et al.~1997), from the extreme-ultraviolet (EUV) imager SOHO/EIT (Berghmans et al.~1998; Krucker and Benz 2000), and from the EUV imager on TRACE (Aschwanden et al.~2000; Aschwanden and Parnell 2002; Aschwanden and Aschwanden 2008a,b). Most of these studies provide statistics on spatial length scales $L$, areas $A$, and durations $T$ of flares (down to nanoflares), but there exists no study to our knowledge that provides statistics on the spatio-temporal evolution $L(t)$ or $A(t)$ of solar flares.  In this paper we are going to analyze the spatio-temporal evolution of all large (GOES X- and M-class) flares observed during the first two years of the SDO mission, which were observed with high spatial resolution ($0.6\\arcsec$), high cadence (12 s), and in 7 coronal wavelengths filters that cover a wide temperature range ($T \\approx 0.5 - 16$ MK). This is an ideal data set for such a study, because both spatial and temporal parameters can be measured with unprecedented accuracy and 100\\% time coverage. Setting a threshold of $>$M1.0 GOES class, we obtain a complete set of 155 flare events detected with both GOES and AIA, which makes it to a perfect representative statistical sample. The data analysis presented here is restricted to the 335 \\ang\\ filter, which appears to be very suitable to capture the high-temperature component of the energy release and heating phase of these largest flares. We fit then various theoretical models to the measured time profiles of the flare areas $A(t)$ or mean radius $r(t) \\propto A(t)^{1/2}$, such as: classical diffusion, $r(t) \\propto t^{1/2}$; sub-diffusion $r(t) \\propto t^{\\beta/2}$ with $\\beta < 1$; super-diffusion or L\\'evy flights, $r(t) \\propto t^{\\beta/2}$ with $\\beta > 1$; or logistic growth, where the time profile initially expands exponentially and then saturates at a finite level, a limit that is also called {\\sl carrying capacity} of limited resources in ecological models. The generic form of these spatio-temporal evolution functions $r(t)$ are shown in Fig.~1. From the fits of the theoretical models to the observed flare data we aim then to gain physical insights into the underlying dynamic processes, which concern the spatial propagation (or chain reaction) of magnetic reconnection or other nonlinear energy dissipation processes. A particular statistical model that we test is the so-called {\\sl fractal-diffusion self-organized criticality model}, which predicts specific spatio-temporal evolutions and powerlaw distributions of the spatial and temporal parameters.  The plan of this paper consists of a brief description of relevant theoretical models (Section 2), the statistical data analysis and forward-fitting to 155 flare events observed with AIA/SDO (Section 3), a discussion of the interpretation and consequences of the results (Section 4), conclusions (Section 5), and a generalization of the FD-SOC model (Appendix A).  \\clearpage ",
        "conclusions": "We analyzed the spatio-temporal evolution in the 155 largest solar flares (M and X-class) observed by GOES and AIA/SDO during the first two years of the SDO mission. We fitted the radial expansion $r(t)$ of flare areas detected in the 335 \\ang\\ filter above some threshold level with a generalized diffusion model that includes the classical diffusion, anomalous diffusion, and the logistic growth limit. The major results and conclusions are:  \\begin{enumerate} \\item{The flare area in all events can be fitted with a radial expansion model $r(t)$ that consists of an initial acceleration phase with exponential growth and a deceleration phase that follows anomalous diffusion, $r(t) \\propto \\kappa t^{\\beta/2}$, with $\\beta=0.53\\pm0.27$ mostly falling into the sub-diffusive regime. The most extreme cases range from logistic growth ($\\beta=0.04$) to super-diffusion ($\\beta=1.35$). The sub-diffusive characteristics is likely to reflect the anisotropic propagation of energy release in a magnetically dominated plasma. The limit of logistic growth indicates the times when the boundaries of energy release regions are reached.}  \\item{The diffusion coefficient $\\kappa = 53\\pm23$ km s$^{-\\delta/2}$, which corresponds to an area diffusion constant of $D \\approx 300-1600$ km$^2$ s$^{-1}$ is found to be significantly faster than diffusion processes measured in the photosphere ($D \\approx 100-600$ km s$^{-1}$) and cannot be explained with crossfield diffusion in the corona, which is strongly inhibited by the low plasma-$\\beta$ parameter.}  \\item{The average diffusion speed during flares is measured in the range of $v \\approx 5-100$ km s$^{-1}$, with maximum speeds of $v_{max} \\approx 10-500$ km s$^{-1}$, which is slower than the sound speed of $c_s \\approx 500-900$ km s$^{-1}$ expected in flares with temperatures of $T_e \\approx 10-30$ MK, but is compatible with the hard X-ray footpoint motion along the neutral line during flares, and thus is likely to represent the mean propagation speed of subsequently triggered magnetic reconnection sites.}  \\item{The fractal-diffusive self-organized criticality model (FD-SOC) describes the diffusive flare progression in a fractal geometry and predicts powerlaw distributions with slopes of $\\alpha_L=3.0$ for length scales $L$, $\\alpha_T=1.1-2.0$ for time scales $T$, and $\\alpha_P=1.83$ for peak fluxes $P$, based on the observed diffusion index range of $\\beta \\approx 0.1-1.0$, the observed flux-volume scaling of $F_{335} \\propto (dV/dt)^{0.8}$, in a 3D-space geometry with fractal dimension $D_3=2.0$. Our observations yield $\\alpha_L=2.0$, $\\alpha_T=2.2$, and $\\alpha_P=1.34$, which is not inconsistent with the FD-SOC model, given the large uncertainties of the powerlaw slopes determined in a small statistical sample. Nevertheless, the evolutionary fits $r(t) \\propto t^{\\beta/2}$ of individual flares confirm the diffusive property of the FD-SOC model for the case of solar flares with a high degree of accuracy, which has been previously predicted based on cellular automaton simulations.} \\end{enumerate}  Thus, the major accomplishment of this study is the experimental proof of a diffusive flare expansion process that has been predicted by the fractal-diffusive self-organized criticality (FD-SOC) model. However, while the classical FD-SOC model assumed classical diffusion ($\\beta=1$), the observational results from AIA/SDO revealed mostly sub-diffusion $(\\beta \\lapprox 1)$, down to the limit of logistic growth ($\\beta \\gapprox 0$). Future studies may extend the statistics of spatio-temporal parameters and allow us to test the predicted size distributions of SOC models with higher accuracy and establish physical scaling laws for flares that are based on the fundamental geometric parameters of space and time."
    },
    "1208/1208.4843_arXiv.txt": {
        "abstract": "In a sample of 54 galaxy clusters $(0.04<z<0.15)$ containing 3551 early-type galaxies suitable for study, we identify those with tidal features both interactively and automatically. We find that $\\sim$ 3\\% have tidal features that can be detected with data that reaches a $3 \\sigma$ sensitivity limit of 26.5 mag arcsec$^{-2}$.  Regardless of the method used to classify tidal features, or the fidelity imposed on such classifications, we find a deficit of tidally disturbed galaxies with decreasing clustercentric radius that is most pronounced inside of $\\sim 0.5R_{200}$.  We cannot distinguish whether the trend arises from an increasing likelihood of recent mergers with increasing clustercentric radius or a decrease in the lifetime of tidal features with decreasing clustercentric radius.  We find no evidence for a relationship between local density and the incidence of tidal features, but our local density measure has large uncertainties.  We find interesting behavior in the rate of tidal features among cluster early-types as a function of clustercentric radius and expect such results to provide constraints on the effect of the cluster environment on the structure of galaxy halos, the build-up of the red sequence of galaxies, and the origin of the intracluster stellar population. ",
        "introduction": "Galaxies and clusters of galaxies assemble hierarchically, through a long progression of mergers and accretion. As one considers galaxy-scale dark matter halos, baryonic physics becomes more integral and so the analytic formalism developed to calculate the rate of dark matter halo evolution \\citep{ps,bond,lc93}, which is broadly confirmed by numerical simulations \\citep{lc94,sheth}, becomes less prescriptive. To remedy this shortcoming, galaxy formation is now studied using hydrodynamic numerical simulations \\citep[e.g.,][]{springel,hoffman}, although the details of the assembly process are sensitive both to baryonic physics below scales accessible in the simulations (``sub-grid\" physics) and to the complex history that occurs on cosmological scales. Tests of such models include comparisons to the properties of galaxy clusters, their constituent galaxies, and the generation of intracluster stars \\citep{puchwein}.  The rate of accretion events and mergers is a fundamental quantity that hierarchical growth models must be able to predict accurately.  A proper comparison to an empirical rate measurement is elusive \\citep[see][for a description of the current state of the field]{lotz}.  Efforts to measure this rate, particularly as a function of redshift, have focused either on measuring close pairs of galaxies \\citep[a few examples include][]{carlberg, patton, lefevre,kartaltepe,deravel}, or identifying galaxies that appear morphologically to be the result of ongoing or recent merger events \\citep[a few examples include][]{abraham,lotz08,jogee,bridge,miskolczi} and results often disagree among studies \\citep{lotz}.  These measurements are difficult, even with the superb angular resolution of the {\\sl Hubble Space Telescope}, for high redshift galaxies.  However, renewed focus on identifying tidal features at low redshift has yielded some striking successes both in the increased incidence of detected features \\citep{tal} and their scale and morphologies \\citep{md}.  We aim to make a measurement of the major merger rate at low redshift, but as a function of environment. A complication in practice is that the definition of a major, vs a minor, merger is typically made on the basis of the mass ratio between the two galaxies. Such a definition is difficult to apply empirically, particularly in environments that may have affected the individual galaxies and in cases where the two initial galaxies have merged into a single remnant. We define a major merger here as one that results in the large-scale tidal features we are searching for around luminous, early-type ($R^{1/4}$ profile) galaxies. Although this approach evidently complicates comparisons to models, it should not impact our primary goal of searching for patterns in the occurrence of tidal features as a function of environment.  \\cite{dubinski} showed how the dark matter halo structure of the merging galaxies can significantly alter the appearance of a major merger, thereby demonstrating that measurements of major merger signatures, even at a single redshift, could be a sensitive model diagnostic.  Beyond the purpose of understanding hierarchical accretion better, we need to study the merger rate in clusters and in the surrounding environments because these events contribute stars to the ubiquitous intracluster stellar population \\citep{zibetti,gonzalez05}, which is an important component in estimates of the baryon budget of groups and clusters \\citep{gonzalez07,giodini,mcgee}, and calculations of the chemical enrichment history of the intracluster medium \\citep{zaritsky,sivanandam}.  Furthermore, extending any study of galaxies within clusters out to the surrounding environs is critical because differences in their star formation histories, which could be related to interactions, begin to appear at clustercentric radii of several Mpc \\citep{lewis,gomez}.  The principal observational challenge in identifying tidal features is that such features are of low surface brightness and have irregular morphologies.  A variety of techniques have been developed to identify merging galaxies and mergers \\citep[e.g.,][]{conselice,lotz08}, but the problem takes on different forms depending on the redshift of the galaxies and the depth and uniformity, or ``flatness\", of the data.  For galaxies in the local universe, \\cite{colbert} employed unsharp masking (subtracting a smoothed image of the galaxy from itself), galaxy model division (dividing the galaxy image by a best-fit model galaxy made from nested elliptical isophotes), and color mapping (division of the galaxy image in one color band by the galaxy image in another color band) to aid in the visual detection of tidal features and found that 41\\% (9 of 22) of analyzed isolated galaxies show tidal features, while only 8\\% (1 of 12) of group galaxies had such features.  While this result is certainly suggestive of lower rates of tidal features in group environments than in the field, Poisson statistics limits the significance of the discrepancy to less than $2\\sigma$. \\cite{vandokkum} divided his galaxy images by a model fit and measured the mean absolute deviation of the residuals to define a quantitative threshold for the identification of tidally disturbed galaxies.  He found that 53\\% of his sample of 126 red, field galaxies show tidal features, while 71\\% of the bulge-dominated early-type subset of 86 galaxies appear to be tidally disturbed. Using a similar method for a sample of nearby, luminous, elliptical galaxies, \\cite{tal} found that 50\\% (5 of 10) of cluster galaxies, 62\\% (13 of 18) of poor-group galaxies, and 76\\% (16 of 21) of isolated galaxies have tidal features.  Similar to the results of \\cite{colbert}, these also suggest that tidal features are less common in cluster environments than in the field, but again the significance of the result is limited by the sample size. \\cite{jan10} quantify the extent of tidal substructure in five elliptical galaxies in the Virgo cluster by measuring the luminosity of the substructure and find no obvious correlation between clustercentric radius and the amount of substructure.  Larger samples are needed to resolve this issue.  Larger samples studied thus far are from field surveys.  \\cite{bridge} visually identify tidal tails in a sample of 27,000 galaxies over 2 square degrees of the Canada-France-Hawaii Telescope Legacy Deep Survey (CFHTLS-Deep) to find that the merger fraction of galaxies evolves from 4\\% at $z \\sim 0.3$ to 19\\% at $z \\sim 1$.  \\cite{miskolczi} study a sample of 474 galaxies selected from the SDSS DR7 archive to find that at least 6\\% of the galaxies have distinct tidal streams and a total of 19\\% show faint features.  No large samples have included substantial numbers of galaxies in dense environments.  Deep, uniform imaging is a prerequisite for any such study. The incidence of identified mergers among field galaxies increases from $<$ 10\\% locally in relatively shallow imaging \\citep{miskolczi} to many tens of percent in deeper imaging \\citep{md, tal}.  For this and other reasons, it is neither surprising nor incorrect that the apparent total rates from different studies vary widely. The key is therefore to compare within a given study rather than across studies.  We identify tidally disturbed galaxies in the largest existing sample of deep, wide-field images of nearby galaxy clusters $(0.04 < z < 0.15)$. These come from the Multi-Epoch Nearby Cluster Survey \\citep[MENeaCS; ][]{sand2}, and this study is part of a series using those data to address a series of science questions including so far the incidence of intracluster supernovae \\citep{sand1}, the rate of SNe Ia in clusters \\citep{sand2}, the evolution of the dwarf-to-giant ratio \\citep{bildfell}, and the rate of SNe II's in clusters \\citep{graham}. The larger sample size enables us to examine, for the first time, relative rates of tidal features as a function of different measures of the global environmental (projected clustercentric radius) and local environment (projected local galaxy density). In \\S2 we briefly present the data we use and discuss our analysis of these data, including the selection of candidate cluster galaxies and the calculation of the tidal parameter. In \\S3 we present our results and discuss whether any significant trends are identified. We present our conclusions in \\S4.  We adopt $H_{0}=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}=0.27$, and $\\Omega_{\\Lambda}=0.73$ for conversions to physical units. ",
        "conclusions": "In a sample of 54 galaxy clusters $(0.04<z<0.15)$ containing 3551 early-type galaxies suitable for study, we identify those with tidal features both interactively and automatically. This constitutes the largest sample studied to date for signs of environmental dependence in the incidence of tidal features. We find tidal features in $\\sim$ 3\\% of galaxies in this sample, with data of this particular depth.  Regardless of the method used to classify tidal features, or the fidelity imposed on such classifications, we find a deficit of tidally disturbed galaxies with decreasing clustercentric radius that is most pronounced inside of $\\sim 0.5R_{200}$.  Although this trend could be attributed to a rise in galaxy-galaxy interactions at large clustercentric distances, where the galaxy pair velocities might be better tuned to mergers, an alternative interpretation is that tidal features are preferentially erased at small clustercentric radii. We examine this hypothesis with a simple toy model that links tidal feature survival time to the local dynamical time and find that although qualitatively the effect is in the correct sense to explain the data, the model falls somewhat short quantitatively. Nevertheless, given the relative success and the limited statistics, we cannot exclude this model as a possible explanation for the radial trend in the incidence of tidal tails.  We also search for a dependence of the incidence of tidal features on local density and find no statistically significant evidence for such, even after accounting for the correlation between clustercentric radius and local density. Unfortunately, the large uncertainties in the measure of local environment limit the interpretation of a null result.  We demonstrate that interesting behavior exists in the rate of tidal features among cluster early-types as a function of clustercentric radius. This measure should provide some guidance for models of merging among early-types, which is conjectured to be significant in the build-up of the red sequence \\citep{bell,faber}.  Models of merging in and around clusters should aim to provide merger estimates based on observable features rather than theoretical constructs. We expect that such work will provide strong constraints on the effect of the cluster environment on the structure of galaxy halos, the build-up of the red sequence of galaxies, and the origin of the intracluster stellar population."
    },
    "1208/1208.3464_arXiv.txt": {
        "abstract": "We have exploited the {\\it Hubble Space Telescope} ({\\it HST}) CANDELS $J$ and $H$-band WFC3/IR imaging to study the properties of (sub)millimetre galaxies within the GOODS-South field. After using the deep radio (VLA 1.4\\,GHz) and {\\it Spitzer} (IRAC 8\\,$\\mu$m) imaging to identify galaxy counterparts for the (sub)millimetre sources, we have then utilised the new CANDELS WFC3/IR imaging in two ways. First, the addition of new deep near-infrared photometry from both {\\it HST} and (at $K$-band) the VLT to the existing GOODS-South database has enabled us to derive improved photometric redshifts and stellar masses, confirming that the (sub)millimetre sources are massive ($\\langle {M}_{\\star} \\rangle = 2.2\\times10^{11} \\pm 0.2\\,{\\rm M_{\\odot}}$) galaxies at $z \\simeq 1-3$. Second, we have exploited the depth and resolution of the WFC3/IR imaging to determine the sizes and morphologies of the galaxies at rest-frame optical wavelengths $\\lambda_{\\rm{rest}} > 4000$\\AA. Specifically, we have fitted two-dimensional axi-symmetric galaxy models to the WFC3/IR images, varying luminosity, axial ratio, half-light radius $r_{1/2}$, and S\\'{e}rsic index $n$. Crucially, the wavelength and depth of the WFC3/IR imaging enables modelling of the mass-dominant galaxy, rather than the blue high surface-brightness features which often dominate optical (rest-frame ultraviolet) images of (sub)millimetre galaxies, and can confuse visual morphological classification. As a result of this analysis we find that $>95$\\% of the rest-frame optical light in almost all of the (sub)millimetre galaxies is well-described by either a single exponential disk ($n \\simeq 1$), or a multiple-component system in which the dominant constituent is disk-like. We demonstrate that this conclusion is completely consistent with the results of recent high-quality ground-based $K$-band imaging sampling even longer rest-frame wavelengths, and explain why. These massive disk galaxies are reasonably extended ($\\langle r_{1/2} \\rangle = 4.5 \\pm 0.5$\\,kpc; median $r_{1/2}=4.0$\\,kpc), consistent with the sizes of other massive star-forming disks at $z \\simeq 2$. In many cases we find evidence of blue clumps within the sources, with the mass-dominant disk component becoming more significant at longer wavelengths. Finally, only a minority of the sources show evidence for a major galaxy-galaxy interaction. Taken together, these results support the view that most (sub)millimetre galaxies at $z \\simeq 2$ are simply the most extreme examples of normal star-forming galaxies at that era. Interestingly, the only two bulge-dominated galaxies are also the two lowest-redshift sources in the sample ($z \\simeq 1$), a result which may reflect the structural evolution of high-mass galaxies in general. ",
        "introduction": "Since (sub)millimetre galaxies were first discovered with the SCUBA camera on the James Clerk Maxwell Telescope in the late 1990s, it has been known that they are violently star-forming galaxies at high redshifts. Indeed, while effective sub-mm/mm imaging took a long time to arrive, the strongly negative k-correction provided by the spectral energy distribution (SED) of dust-enshrouded star formation (i.e. a modified black body -- e.g. Hughes et al. 1993; Blain et al. 2002) resulted in the first sub-mm galaxies being discovered at rather high redshifts (i.e. $z \\simeq 2 - 4$; Hughes et al. 1998). Interestingly, this meant that the supporting multi-frequency data available at the time was of insufficient quality to adequately reveal the true nature of these sources, and it has taken over a decade of effort with ever-improving radio surveys, optical-infrared imaging/spectroscopy with ground-based 8-10\\,m telescopes, and deep {\\it Spitzer} and {\\it Hubble Space Telescope} ({\\it HST}) imaging to begin to clarify the physical nature of (sub)millimetre galaxies, and their role in the process of galaxy formation.  This is still a work in progress. For example, it is only very recently that a spectroscopic redshift was finally measured for the brightest sub-mm galaxy discovered towards the HDF in the first unbiased sub-mm survey at 850\\,$\\mu$m (HDF850.1; Hughes et al. 1998; Downes et al. 1999; Dunlop et al. 2004; Cowie et al. 2009) revealing it to lie at $z = 5.183$ (Walter et al. 2012). This measurement, based on CO lines, has been made possible by recent advances in sub-mm/mm spectroscopy (e.g. Riechers et al. 2010, 2012, and Hodge et al. 2012 for morphology), a technique which should soon enable complete spectroscopic redshift distributions to be established for the (now substantial) samples of (sub)millimetre sources which have been uncovered by continuum surveys with SCUBA, MAMBO, LABOCA and AzTEC (e.g. Coppin et al. 2006; Greve et al. 2008; Weiss et al. 2009; Austermann et al. 2010). With the advent of the Atacama Large Millimetre Array (ALMA), high-resolution sub-mm/mm imaging is set to become straightforward, and detailed sub-mm/mm spectroscopy will soon be routine.  However, as we now enter a new era in sub-mm/mm observational astronomy, it is important to note that major advances in our knowledge of (sub)millimetre galaxies have already been gained from studies at optical-mid-infrared wavelengths. For example, aided by the improved positional accuracy provided by deep VLA 1.4-GHz imaging, it has already proved possible to identify optical/near-infrared counterparts for $\\simeq 80$\\% of the sub-mm/mm galaxies in bright samples, and hence to determine the basic form of their redshift distribution. These studies indicate that (sub)millimetre galaxies peak in number density around $z \\simeq 2$ (e.g. Chapman et al. 2005; Wardlow et al. 2011; Micha{\\l}owski et al. 2012b; Schael et al. 2013), although there is some evidence that the brightest (sub)millimetre galaxies lie predominantly at $z > 3$ (e.g. Smolcic et al. 2012; Koprowski et al. 2013), and the full effects of large-scale structure have yet to be properly clarified (e.g. Micha{\\l}owski et al. 2012b; Scott et al. 2012; Shimizu, Yoshida \\& Okamoto 2012).  Moreover, while one might argue that redshift measurements for (sub)millimetre sources may increasingly be made at sub-mm/mm wavelengths, there are some crucial properties of (sub)millimetre-selected galaxies which can {\\it only} be determined from deep optical--mid-infrared data. These include their stellar masses, formation history, and the spatial distribution of their existing stellar populations. Measurement of these physical quantities is of crucial importance if we are to understand the nature of (sub)millimetre galaxies as a function of luminosity and redshift, and place them in the context of the general galaxy population at comparable redshift/stellar-mass. Importantly, this is now an achievable goal, as the near/mid-infrared study of more moderate star-forming galaxies has now reached out to comparable redshifts (e.g. Daddi et al. 2007; Elbaz et al. 2011) and the construction of significant and complete mass-selected samples down to $M_{\\star} \\simeq 5 \\times 10^{10}\\,{\\rm M_{\\odot}}$ is now possible out to $z \\simeq 3$ (e.g. Bruce et al. 2012).  The most important recent advance in this context is the provision of deep, high-quality, near-infrared imaging provided by WFC3/IR on {\\it HST}. In particular, the CANDELS survey (Grogin et al. 2011; Koekemoer et al. 2011) is providing homogeneous WFC3/IR $J_{125}$ and $H_{160}$ imaging of five fields which overlap significantly with existing sub-mm/mm surveys. Such imaging is of general importance for the study of galaxies at $z \\simeq 1-3$, as it offers high quality {\\it rest-frame optical} imaging at these redshifts where, until recently, our knowledge of galaxy morphologies and colours was largely based on the rest-frame ultraviolet. However, it is arguably of special importance for (sub)millimetre galaxies, where much of the rest-frame UV light may be obscured by dust, and hence morphological k-corrections can be potentially dramatic.  In this paper we exploit the CANDELS data in the GOODS-South field to attempt to improve our knowledge of the physical properties of (sub)millimetre galaxies. Of the five CANDELS fields, GOODS-South currently offers the best overlap between WFC3/IR imaging and existing samples of sub-mm/mm-selected sources (although this situation will soon change as a result of the SCUBA2 Cosmology Redshift Survey\\footnote{www.jach.hawaii.edu/JCMT/surveys/Cosmology.html}), with the added advantages of very deep {\\it Spitzer}, {\\it HST} ACS, and VLT imaging, and high-density spectroscopy (e.g. Vanzella et al. 2008). Our aim was to utilise the WFC3/IR imaging (Grogin et al. 2011) to first improve the measurement of the photometric redshifts and stellar masses of the known LABOCA/AzTEC sources in this field (Weiss et al. 2009; Scott et al. 2010; Wardlow et al. 2011; Yun et al. 2012), and then to determine their rest-frame optical morphologies. Crucially, because of the complete, homogeneous, and unbiased nature of the CANDELS imaging, we have also been able to compare the rest-frame optical morphologies of the (sub)millimetre galaxies to those displayed by the general population of galaxies at comparable redshift and stellar mass (Bruce et al. 2012).  This latter point is one of the key advantages offered by CANDELS, alongside, of course, the major improvement in depth and angular-resolution offered by WFC3/IR over previous near-infrared imagers. To date, deep, near-infrared imaging has only been obtained for a small number of (sub)millimetre galaxies (e.g. Smail et al. 2003), and generally with inadequate resolution for the reliable extraction of basic morphological parameters. In an attempt to exploit the resolution advantage offered by {\\it HST}, Swinbank et al. (2010) reported the results of {\\it HST} NICMOS+ACS imaging of a sub-sample of (sub)millimetre galaxies. While S\\'{e}rsic indices for individual galaxies were not reported, median values of $n_{i}=1.8 \\pm 1.0$ and $n_{H}=1.4 \\pm 0.8$ were given, and additional measurements of S\\'{e}rsic index based on a stacked NICMOS imaging found $n_{H}=2.0\\pm0.5$. Combining these results with measurements of half-light radii and asymmetry, Swinbank et al. (2010) suggested that the stellar structure of (sub)millimetre galaxies was best described by a spheroid/elliptical galaxy light distribution, despite the low median S\\'{e}rsic index. More recently, Targett et al. (2011) analysed deep, high-quality ground-based $K$-band images of complete subsamples of (sub)millimetre galaxies at redshifts $z \\simeq 2-3$. While the use of $K$-band ensures imaging longward of the 4000\\AA/Balmer break out to $z \\simeq 4$, the ground-based images have poorer resolution than {\\it HST} data. Nevertheless, Targett et al. (2011) found the galaxies they studied to be reasonably extended ($\\langle r_{1/2} \\rangle = 3.4 \\pm 0.3$\\,kpc; median $r_{1/2}=3.1$\\,kpc) exponential disks ($\\langle n \\rangle = 1.44 \\pm 0.16$; median $n = 1.08$). Thus, while there is improving agreement over the basic sizes of (sub)millimetre galaxies, there remains controversy over their morphologies.  There is also an ongoing debate over the stellar masses of (sub)millimetre galaxies, and whether or not their extreme star-formation rates are triggered predominately by major galaxy-galaxy mergers. As discussed in Dunlop (2011) and Micha{\\l}owski et al. (2012a), this strikes to the heart of the nature of (sub)millimetre galaxies, with important implications for refining current models of galaxy formation. For a long time it was believed that (sub)millimetre galaxies were the consequence of major galaxy mergers, in part because the brightest ULIRGs uncovered by IRAS in the local Universe appeared to be merger driven (e.g. Sanders \\& Mirabel 1996). However, over recent years it has become clear that the specific star-formation rate ($sSFR$=star-formation rate/stellar mass) of the general star-forming galaxy population is a factor $\\simeq 20$ higher at $z \\simeq 2$ than at the present day (Daddi et al. 2007; Elbaz et al. 2011; Karim et al. 2011), with $sSFR \\simeq 2-4\\,{\\rm Gyr^{-1}}$ (Gonzalez et al. 2010). Consequently, provided they have stellar masses $M_{\\star} > 10^{11}\\,{\\rm M_{\\odot}}$, it can be argued that the extreme star-formation rates inferred for (sub)millimetre galaxies are not ``unexpected'' at $z \\simeq 2$, and that the situation has been confused by their relatively early discovery. This realization has focussed increased attention on the accurate determination of the stellar masses of (sub)millimetre galaxies (e.g. Dye et al. 2008; Micha{\\l}owski et al. 2010a,b; Hainline et al. 2011; Micha{\\l}owski et al. 2012a) and the evidence for galaxy interactions (e.g. Tacconi et al. 2008; Engel et al. 2010; Alaghband-Zadeh et al. 2012).  Importantly, some theoretical models appear to require galaxy interactions, and even the additional boost provided by a top-heavy stellar initial mass function (IMF) to produce sufficient numbers of high-redshift (sub)millimetre sources (Baugh et al. 2005), but this may simply reflect an underlying inability of some semi-analytic models to produce adequate numbers of high-mass galaxies at $z \\simeq 2-3$ (Swinbank et al. 2008). Models involving AGN feedback have improved in this respect (e.g. Bower et al. 2006; Croton et al. 2006), and indeed some authors have argued that the (sub)millimetre galaxy population can be reconciled with the high-mass end of the ``main sequence'' of galaxy formation at $z \\simeq 2 - 3$, and are fed by smooth infall rather than major mergers (e.g. Dav\\'{e} et al. 2010, and see also Finlator et al. 2006; Fardal et al. 2007; Dekel et al. 2009). Moreover, it has been argued that major mergers are not sufficiently common in the redshift range of interest to explain the (sub)millimetre population, and that in any case the ability of galaxy mergers to significantly enhance the already high specific star-formation rates of galaxies at $z \\simeq 2 - 3$ is relatively modest, $\\simeq 10-25$\\% (e.g. Cen 2012, Kaviraj et al. 2012). Of course the true situation may be complex, and recently arguments have been advanced that the (sub)millimetre galaxy population may contain a mix of both massive ``normal'' star-forming galaxies and a subset of extreme objects boosted by interactions (e.g. Hayward et al. 2012).  Clearly, therefore, further observational clarification of the masses and morphologies of (sub)millimetre galaxies is of value, especially if the results can be robustly compared with the properties of the general galaxy population. This is facilitated here by the work of Bruce et al. (2012), who utilised comparable CANDELS imaging to study the rest-frame optical morphologies of $\\simeq 200$ massive ($M_{\\star}\\ge10^{11}\\,{\\rm M_{\\odot}}$) galaxies in the the redshift range $1<z<3$. This work has shown that the CANDELS ``Wide'' data are of sufficient quality and depth to enable the robust measurement of the morphologies of such massive galaxies out to $z \\simeq 3$, and has revealed the emergence of a predominantly disk-dominated population by $z \\ge 2$. This pattern is confirmed at somewhat lower stellar masses by Law et al. (2012a) who, from an {\\it HST} study of 306 star-forming galaxies at $1.5 < z < 3.6$ with $M_{\\star} = 10^{9}-10^{11}\\,{\\rm M_{\\odot}}$, found that star-forming galaxies at these redshifts are best described by an exponential (disky) surface brightness profile. However, while there appears to be a growing consensus that massive galaxies at $z > 2$ are disk-like, there is an interesting debate over how these high-redshift disks compare to their local counterparts, since they appear to display a peaked axial-ratio distribution (e.g. Law et al. 2012a; Bruce et al. 2012), and have substantially higher velocity dispersions than low-redshift disk galaxies (e.g. F\\\"{o}rster-Schreiber et al. 2009; Law et al. 2012b). There is also considerable debate over the prevalence and potential importance (or otherwise) of clumps in these high-redshift star-forming disks (e.g. Wuyts et al. 2012), and the extent to which they are predicted by current models of galaxy formation (e.g. Ceverino et al. 2012; Genel et al. 2012). Whatever their origin and importance, it is clear that such high-surface brightness, and generally blue (e.g. Guo et al. 2012) clumps mitigate against the effective morphological classification of high-redshift star-forming galaxies from optical imaging (which samples the rest-frame near-ultraviolet at $z > 1.5$).  It is clearly important to establish where (sub)millimetre galaxies fit into this picture. For example, as star-forming galaxies they might be expected to be extended disks, and yet, based on their high stellar masses it is hard to avoid the conclusion that they are the progenitors of present-day giant elliptical galaxies. Here we attempt to move this subject forward using the new WFC3/IR CANDELS imaging to study the properties of the AzTEC and LABOCA-selected galaxies in the GOODS-South field. For the first time, we can study the rest-frame optical morphologies of (sub)millimetre galaxies in detail, due to the combination of high angular-resolution (0.1\\,arcsec) and depth ($H_{160} = 26.5$; AB mag, 5-$\\sigma$) provided by the CANDELS WFC3/IR imaging. In addition, the existing rich multi-frequency data in this well-established field provides a significant number of spectroscopic redshifts, and enables the most accurate photometric redshifts achieved for (sub)millimetre galaxies to date (with consequently reduced uncertainties in derived stellar masses). We note that it is important to distinguish between the evidence for or against merger-driven starbursts and the nature of the dominant underlying galaxy; in this paper our focus is on determining the latter and comparing the stellar masses and morphologies of (sub)millimetre galaxies to the properties of the general galaxy population at $z \\simeq 2 -3$. We do also consider the imaging evidence for major galaxy-galaxy interactions, but such work is arguably better pursued with the velocity information gleaned from optical and/or mm spectroscopy. It does, however, need to be performed in an unbiased manner including both (sub)millimetre bright and non (sub)millimetre-detected massive galaxies at $z \\simeq 2$, before definitive conclusions can be drawn regarding the importance of major mergers in producing the (sub)millimetre galaxy population.  This paper is structured as follows. In Section 2 we summarise the available (sub)millimetre survey data within GOODS-South, the new CANDELS {\\it HST} imaging, and the additional key supporting data available within this field. Next, in Section 3, we explain how we determined the radio and IRAC IDs, and hence WFC3/IR $H$-band galaxy counterparts for the (sub)millimetre sources. Then, in Section 4 we describe the two-dimensional modelling used to extract the basic morphological properties of the galaxies, and the SED fitting to the multi-frequency data used to determine redshifts and stellar masses. In Section 5 we present the results of our analysis, and also place our findings in the context of studies of other galaxy populations at both high and low redshift. Our main conclusions are summarised in Section 6. Throughout we quote magnitudes in the AB system (Oke 1974), and assume a cosmological model with $H_0 = 70\\,{\\rm km s^{-1} Mpc^{-1}}$, $\\Omega_{\\Lambda} = 0.7$, and $\\Omega_m = 0.3$.   \\begin{table*} \\begin{center} \\caption{Positional information on the 24-source combined AzTEC+LABOCA GOODS-South (sub)millimetre galaxy sample utilised in this paper, and the adopted galaxy counterparts in the {\\it HST} CANDELS $H_{160}$ imaging. The first column gives the name of the AzTEC source from Scott et al. (2010), with the next two columns giving the AzTEC position (AzTEC.GS19 is tabulated twice here because it has two alternative IDs, as explained in Section 3). The names of sources with statistically secure ($P < 0.1$) counterparts are highlighted in bold. Column 4 then gives the source name from the LESS survey (Weiss et al. 2009), where here the name itself contains the original LABOCA survey position. Column 5 gives the de-boosted flux densities of the (sub)milimetre sources at 1.1\\,mm (AzTEC) and 870\\,${\\rm \\mu m}$ (LABOCA). Columns 6 and 7 then give the positions of the $H_{160}$ galaxy counterpart we have selected for morphological study. These galaxy identifications were selected as described in Section 3. Column 8 then gives the probability that the adopted galaxy counterpart is merely a chance association, while column 9 indicates whether the galaxy identification was selected on the basis of the radio or {\\it Spitzer} imaging, and the source of the quoted $P$ value; (a) our identification process, or (b) from Yun et al. (2012).} \\begin{tabular}{lllllllll} \\hline\\hline AzTEC             & RA          & Dec.          & LABOCA                    & Flux        & CANDELS      & CANDELS      & $P$   & Type\\\\ ID                & (J2000)     & (J2000)       & ID                        & (mJy)       & RA  (J2000)  & Dec (J2000)  &       &\\\\ \\hline \\bf{AzTEC.GS03}   & 03:32:47.86 & $-$27:54:19.3 & \\bf{LESSJ033248.1-275414} & 4.8$\\pm$0.6 & 03:32:47.992 & -27:54:16.42 & 0.017 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS06}   & 03:32:25.73 & $-$27:52:19.4 & \\bf{LESSJ033225.7-275228} & 3.6$\\pm$0.5 & 03:32:25.258 & -27:52:30.51 & 0.062 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS08}   & 03:32:05.12 & $-$27:46:45.8 & \\bf{LESSJ033205.1-274652} & 3.4$\\pm$0.6 & 03:32:04.873 & -27:46:47.37 & 0.015 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS11}   & 03:32:15.79 & $-$27:50:36.8 &                           & 3.3$\\pm$0.6 & 03:32:15.319 & -27:50:37.12 & 0.067 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS12}   & 03:32:29.13 & $-$27:56:13.8 & \\bf{LESSJ033229.3-275619} & 5.1$\\pm$1.4 & 03:32:29.290 & -27:56:19.47 & 0.032 & b - 1.4\\,GHz\\\\ \\bf{AzTEC.GS13}   & 03:32:11.91 & $-$27:46:16.9 &                           & 3.1$\\pm$0.6 & 03:32:11.908 & -27:46:15.51 & 0.009 & b - 1.4\\,GHz\\\\ \\bf{AzTEC.GS16}   & 03:32:37.67 & $-$27:44:01.8 &                           & 2.7$\\pm$0.5 & 03:32:38.006 & -27:44:00.57 & 0.066 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS17}   & 03:32:22.31 & $-$27:48:16.4 &                           & 2.9$\\pm$0.6 & 03:32:22.566 & -27:48:14.83 & 0.026 & a/b - 8/24\\,${\\rm \\mu m}$\\\\ \\bf{AzTEC.GS18}   & 03:32:43.58 & $-$27:46:36.9 & \\bf{LESSJ033243.6-274644} & 3.1$\\pm$0.6 & 03:32:43.536 & -27:46:38.90 & 0.035 & b - 1.4\\,GHz\\\\ \\bf{AzTEC.GS19-1} & 03:32:23.21 & $-$27:41:28.8 &                           & 2.6$\\pm$0.5 & 03:32:22.879 & -27:41:24.96 & 0.074 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS19-2} & 03:32:23.21 & $-$27:41:28.8 &                           & 2.6$\\pm$0.5 & 03:32:22.708 & -27:41:26.39 & 0.088 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS21}   & 03:32:47.60 & $-$27:44:49.3 &                           & 2.7$\\pm$0.6 & 03:32:47.592 & -27:44:52.26 & 0.026 & a - 1.4\\,GHz\\\\ AzTEC.GS22        & 03:32:12.60 & $-$27:42:57.9 &                           & 2.1$\\pm$0.6 & 03:32:12.545 & -27:43:05.99 & 0.116 & b - 1.4\\,GHz\\\\ \\bf{AzTEC.GS23}   & 03:32:21.37 & $-$27:56:28.1 & \\bf{LESSJ033221.3-275623} & 4.7$\\pm$1.4 & 03:32:21.574 & -27:56:23.94 & 0.054 & a - 1.4\\,GHz\\\\ AzTEC.GS24        & 03:32:34.76 & $-$27:49:43.1 &                           & 2.3$\\pm$0.6 & 03:32:34.163 & -27:49:39.52 & 0.141 & b - 1.4\\,GHz\\\\ AzTEC.GS26        & 03:32:15.79 & $-$27:43:36.6 &                           & 2.2$\\pm$0.5 & 03:32:16.281 & -27:43:43.41 & 0.624 & a - 8\\,${\\rm \\mu m}$\\\\ AzTEC.GS27        & 03:32:42.42 & $-$27:41:51.9 &                           & 2.2$\\pm$0.6 & 03:32:41.632 & -27:41:51.45 & 0.623 & a - 8\\,${\\rm \\mu m}$\\\\ AzTEC.GS28        & 03:32:42.71 & $-$27:52:06.8 &                           & 2.1$\\pm$0.5 & 03:32:41.853 & -27:52:02.47 & 0.734 & a - 8\\,${\\rm \\mu m}$\\\\ \\bf{AzTEC.GS30}   & 03:32:20.94 & $-$27:42:40.8 &                           & 1.8$\\pm$0.5 & 03:32:20.658 & -27:42:34.40 & 0.082 & b - 1.4\\,GHz\\\\ AzTEC.GS34        & 03:32:29.77 & $-$27:43:13.1 &                           & 1.7$\\pm$0.5 & 03:32:29.479 & -27:43:22.10 & 0.171 & a - 1.4\\,GHz\\\\ \\bf{AzTEC.GS35}   & 03:32:26.90 & $-$27:40:52.1 &                           & 2.1$\\pm$0.6 & 03:32:27.181 & -27:40:51.42 & 0.008 & b - 1.4\\,GHz\\\\ \\bf{AzTEC.GS38}   & 03:32:09.26 & $-$27:42:45.5 &                           & 1.7$\\pm$0.6 & 03:32:09.705 & -27:42:48.10 & 0.015 & a - 1.4\\,GHz\\\\ &             &               & \\bf{LESSJ033217.6-275230} & 6.3$\\pm$1.3 & 03:32:17.617 & -27:52:28.52 & 0.016 & a - 8\\,${\\rm \\mu m}$\\\\ &             &               & \\bf{LESSJ033219.0-275219} & 9.1$\\pm$1.2 & 03:32:19.032 & -27:52:13.97 & 0.001 & b - 1.4\\,GHz\\\\ &             &               & \\bf{LESSJ033243.3-275517} & 5.2$\\pm$1.4 & 03:32:43.179 & -27:55:14.49 & 0.019 & a - 1.4\\,GHz\\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "We have exploited the {\\it Hubble Space Telescope} CANDELS $J$ and $H$-band WFC3/IR imaging to study the properties of (sub)millimetre galaxies in the GOODS-South field which have been revealed by the AzTEC 1.1\\,mm and LABOCA 870\\,$\\mu$m surveys of the Extended Chandra Deep Field South). After using the deep radio (VLA 1.4\\,GHz) and {\\it Spitzer} (IRAC 8\\,$\\mu$m) imaging to identify galaxy counterparts for the (sub)millimetre galaxies, we have then utilised the new CANDELS WFC3/IR imaging in two ways.  First, the addition of new deep near-infrared photometry from both {\\it HST} and (at $K$-band) HAWK-I on the VLT to the existing extensive GOODS-South multi-frequency database has enabled us to derive improved photometric redshifts and stellar masses for all the (sub)millimetre sources. Our results confirm that the (sub)millimetre sources are massive ($\\langle {M}_{\\star} \\rangle = 2.2\\times10^{11} \\pm 0.2\\,{\\rm M_{\\odot}}$) galaxies at $z \\simeq 1-3$.  Second, we have exploited the depth and superior angular resolution of the WFC3/IR $H_{160}$ imaging to determine the sizes and morphologies of the galaxies at rest-frame optical wavelengths $\\lambda_{\\rm{rest}} > 4000$\\AA. Crucially, the wavelength and depth of the WFC3/IR imaging enables modelling of the mass-dominant galaxy, rather than the blue high-surface brightness features which often dominate optical (rest-frame ultraviolet) images of sub-mm galaxies, and can confuse visual morphological classification. As a result of this analysis we find that $>95$\\% of the rest-frame optical light in almost all of the (sub)millimetre galaxies is well-described by either a single exponential disk ($n \\simeq 1$), or a multiple-component system in which the dominant constituent is disk-like. We demonstrate that this conclusion is completely consistent with the results of recent high-quality ground-based $K$-band imaging studies of (sub)millimetre galaxies sampling even longer rest-frame wavelengths (Targett et al. 2011), and explain why. We also briefly compare our findings with the results of recent morphological studies of high-redshift ULIRGs selected at much shorter ($\\simeq 10\\times$) far-infrared wavelengths (Kartaltepe et al. 2012).  The massive disk galaxies which host strong (sub)millimetre emission are reasonably extended ($\\langle r_{1/2} \\rangle = 4.5 \\pm 0.5$\\,kpc; median $r_{1/2}=4.0$\\,kpc), consistent with the sizes of other massive star-forming disks at $z \\simeq 2$. In many cases we find evidence of blue clumps within the sources, with the mass-dominant disk component becoming more significant at longer wavelengths. Interestingly, the only two bulge-dominated sub-mm galaxies are also the two lowest-redshift sources in the sample ($z \\simeq 1$), a result which may reflect the structural evolution of high-mass galaxies in general.  Our main result that (sub)millimetre galaxies are almost universally massive, star-forming disk galaxies does not exclude the possibility that some fraction of these objects have their star-formation boosted by interactions/mergers. However, our imaging data indicate that only a minority ($< 25$\\%) of the (sub)millimetre sources in the GOODS-South sample could be regarded as involved in a major galaxy-galaxy interaction. Taken together, our results support the view that most sub-mm galaxies at $z \\simeq 2$ are simply the most extreme examples of normal star-forming galaxies at that era."
    },
    "1208/1208.6288_arXiv.txt": {
        "abstract": "{GRB-selected galaxies are broadly known to be faint, blue, young, star-forming dwarf galaxies. This insight, however, is based in part on heterogeneous samples of optically selected, lower-redshift galaxies. To study the statistical properties of GRB-selected galaxies we here introduce The Optically Unbiased GRB Host (TOUGH) complete sample of 69 X-ray selected \\swift\\ GRB host galaxies spanning the redshift range 0.03--6.30 and summarise the first results of a large observational survey of these galaxies.}  \\FullConference{Gamma-Ray Bursts 2012 Conference -GRB2012,\\\\ May 07-11, 2012\\\\ Munich, Germany}   \\begin{document} ",
        "introduction": "Gamma-ray bursts (GRBs) arise from the deaths of short-lived stars \\citep{2003Natur.423..847H}. Hence, their host galaxies bear important information about the nature of GRB progenitors and furthermore act as tracers of star formation at a very wide range of redshifts.  The host galaxy of the nearest known GRB host galaxy, that of GRB 980425/SN 1998bw at $z=0.0085$ \\citep{1998Natur.395..670G}, may be seen as representative of this population: it is a dwarf (SMC sized) star-forming galaxy harbouring a highly star-forming region \\citep{2000ApJ...542L..89F,2008ApJ...672..817M}. Indeed, in the {\\it BeppoSAX/HETE-II} era, it was established that GRB host galaxies are typically sub-luminous, blue, young, high specific star-formation rate systems \\citep{1999ApJ...519L..13F,2003A&A...400..499L,2004A&A...425..913C}, with the GRBs originating from the UV/blue light of the galaxies \\citep{2002AJ....123.1111B,2006Natur.441..463F}. The metallicities were found to be preferentially low \\citep{2006AcA....56..333S,2008AJ....135.1136M}, providing an interesting constraint on progenitor models of GRBs, consistent, according to the mass-metallicity relation, with the finding that the stellar masses are generally low \\citep{2010ApJ...721.1919C,2009ApJ...691..182S}. Early sub-mm (SCUBA) observations revealed a rare population of low-redshift ($z\\sim 1$) blue galaxies \\citep{2003ApJ...588...99B,2004MNRAS.352.1073T}, with somewhat elevated dust temperatures \\citep{2009ApJ...693..347M}.  The above results are based on heterogeneous samples of mostly lower-redshift galaxies, e.g., relying on the existence of a bright optical afterglow for localisation and redshift determination.  At higher redshifts, GRB host galaxies are very faint, consistent with their being dwarf galaxies. For example, at $z=3.2$ (the redshift being determined via afterglow absorption spectroscopy) GRB 020124 \\citep{2002ApJ...581..981B,2003ApJ...597..699H} and GRB 060526 \\citep{2010A&A...523A..70T} were found to be hosted by galaxies with $R>29.5$ and $R>28.5$, respectively. Remarkably, there are still no high-redshift ($z>5$) GRB host galaxies robustly detected in emission \\citep{2012arXiv1201.6074T}, consistent with their being sub-$L^*$ galaxies.  However, more recent results have demonstrated that the host galaxies of dark GRBs are more chemically evolved and have higher masses \\citep{2011A&A...534A.108K}, complicating the early simple picture of GRB host galaxies being sub-luminous, low-metallicity, highly star-forming blue galaxies (see also the contributions by Kr\\\"uhler, Levesque, and Perley to these proceedings). ",
        "conclusions": ""
    },
    "1208/1208.0184_arXiv.txt": {
        "abstract": "Recent cosmological observations indicate the existence of extra light species, i.e., dark radiation.  In this paper we show that signatures of the dark radiation are imprinted in the spectrum of inflationary gravitational waves. If the dark radiation is produced by the decay of a massive particle, high frequency mode of the gravitational waves are suppressed.  In addition, due to the effect of the anisotropic stress caused by the dark radiation, a dip in the gravitational wave spectrum may show up at the frequency which enters the horizon at the time of the dark radiation production. Once the gravitational wave spectrum is experimentally studied in detail, we can infer the information on how and when the dark radiation was produced in the Universe. ",
        "introduction": "Recently, there are increasing evidence of the extra non-interacting relativistic degrees of freedom, in addition to the standard three (nearly) massless neutrino species.  The abundance of relativistic component is parametrized by the effective number of neutrino species, $N_{\\rm eff}$, as \\begin{equation} \\rho_{\\rm rel}=\\left[ 1+N_{\\rm eff}\\frac{7}{8}\\left(\\frac{4}{11}\\right)^{4/3} \\right] \\rho_\\gamma, \\end{equation} where $\\rho_{\\rm rel}$ is the total relativistic energy density and $\\rho_\\gamma = (\\pi^2/15)T_\\gamma^4$ denotes the photon energy density measured after the $e^+e^-$ annihilation with $T_\\gamma$ representing the photon temperature.  The standard model predicts $N_{\\rm eff}=3.05$~\\cite{Mangano:2005cc}.  The $N_{\\rm eff}$ can be constrained from various observations. First, increasing $N_{\\rm eff}$ leads to larger Hubble expansion rate at the big bang nucleosynthesis epoch, which in turn results in increase of the primordial helium abundance.  Recent observations suggest $N_{\\rm eff} = 3.68^{+0.80}_{-0.70}$ at $2\\sigma$ level~\\cite{Izotov:2010ca}.  (See, however, also Ref.~\\cite{Aver:2010wq} for discussion on the error estimation in the helium abundance.)  The cosmic microwave background (CMB) anisotropy is also sensitive to $N_{\\rm eff}$.  The information on $N_{\\rm eff}$ is imprinted in the CMB anisotropy in some ways.  First, increase of $N_{\\rm eff}$ makes the early integrated Sachs-Wolfe effect more efficient, and the first peak of the CMB power spectrum is enhanced.  Second, it tends to make the scale of sound horizon smaller at the recombination epoch, resulting in shift of the peak positions in the CMB power spectrum toward high multipole moment.  Third, it erases the small scale power spectrum due to the effect of free-streaming.  The WMAP seven-year results combined with standard rulers give $N_{\\rm eff} = 4.32^{+0.86}_{-0.88}$ at $1\\sigma$ level~\\cite{Komatsu:2010fb}. Adding small scale CMB measurements improves the accuracy as $N_{\\rm eff} = 4.56\\pm 0.75$ for ACT~\\cite{Dunkley:2010ge} and $N_{\\rm eff} = 3.86\\pm 0.42$ for SPT~\\cite{Keisler:2011aw} both at $1\\sigma$ level. Ref.~\\cite{Archidiacono:2011gq} combined the WMAP, ACT and SPT datasets with standard rulers and obtained $N_{\\rm eff} = 4.08^{+0.71}_{-0.68}$ at $2\\sigma$ level.\\footnote{ Note that the statistical significance depends on the prior on the Hubble parameter~\\cite{Calabrese:2012vf}. } See also recent related studies~\\cite{Hamann:2011hu,Nollett:2011aa,Hamann:2010bk,Hamann:2011ge}.  To summarize, at the current situation, observations suggest $\\Delta N_{\\rm eff}\\equiv N_{\\rm eff}-3 \\simeq 1$ at nearly $2\\sigma$ level. Motivated by these increasing evidence of the extra light species, which is often called ``dark radiation'', models to explain dark radiation were proposed~\\cite{Ichikawa:2007jv,Jaeckel:2008fi,Nakayama:2010vs,Fischler:2010xz,Kawasaki:2011ym,Hall:2011zq,Hasenkamp:2011em,Kawasaki:2011rc,Menestrina:2011mz,Kobayashi:2011hp,Hooper:2011aj,Jeong:2012hp,Kawakami:2012ke,Blennow:2012de,Moroi:2012vu}. Although there are many candidates, if the dark radiation has only extremely weak interaction with the standard model particles, it may be difficult to detect it experimentally.  Thus it is important to study how to confirm and distinguish models of dark radiation by other observations.  For example, in Refs.~\\cite{Kawasaki:2011rc,Kawakami:2012ke} the possibility that the dark radiation has (non-Gaussian) isocurvature perturbations was considered. In Ref.~\\cite{Zhao:2009aj} the effect of dark radiation on CMB B-mode spectrum is discussed.  In this paper we consider a novel method to detect dark radiation through inflationary gravitational waves (GWs).  It is known that relativistic free streaming fluid can contribute to the anisotropic stress, which potentially affects the propagation of GWs~\\cite{Weinberg:2003ur}.  This effect was concretely studied for the free-streaming neutrinos.  GWs entering the horizon after the neutrino freezeout dissipate their energies and, as a result, a modulation feature shows up in the GW spectrum~\\cite{Weinberg:2003ur,Dicus:2005rh,Watanabe:2006qe}.  It was also studied in the context of large lepton asymmetry~\\cite{Ichiki:2006rn}.  Therefore, it is expected that the dark radiation also induces similar effects on GWs.  As opposed to the case of neutrinos, we do not know when and how the dark radiation was generated in the Universe.  Thus the position and strength of the modulation in the GWs, if detected, tells us exactly about the production mechanism of dark radiation.  In particular, models of dark radiation produced by decay of non-relativistic fields \\cite{Ichikawa:2007jv,Fischler:2010xz,Kawasaki:2011ym,Hasenkamp:2011em,Kawasaki:2011rc,Menestrina:2011mz,Kobayashi:2011hp,Hooper:2011aj,Jeong:2012hp} shows characteristic features in the primordial GW spectrum.  The feature consists of combination of the anisotropic stress effect and the modified background expansion history.  This is detectable in future space-based GW detectors such as DECIGO~\\cite{Seto:2001qf} and BBO~\\cite{Crowder:2005nr,Cutler:2009qv}.  We also show that the GW spectrum will be a powerful tool for confirming the dark radiation produced thermally and decoupled at some epoch in the early Universe.   This paper is organized as follows.  In Sec.~\\ref{sec:DR} we review a model of dark radiation produced by decaying particles.  In Sec.~\\ref{sec:GW} we calculate the evolution of gravitational waves in the presence of anisotropic stress induced by dark radiation, and show that characteristic signatures appear in the spectrum. Sec.~\\ref{sec:conc} is devoted to conclusions and discussion. ",
        "conclusions": "\\label{sec:conc}   In this paper we have studied the spectrum of inflationary GW background in the presence of dark radiation, motivated by recent observational preferences for $\\Delta N_{\\rm eff}\\sim 1$. We have assumed that the dark radiation is non-thermally produced by decay of massive particles $\\phi$. There are several effects on the GW spectrum. First, the equation of state of the Universe is modified due to the $\\phi$ energy density and it changes the shape of the GW spectrum. Second, the anisotropic stress carried by dark radiation dissipates the GW amplitude for modes entering the horizon around and after $\\phi$ decay. Numerical results show that there may appear a characteristic dip around $k\\sim k_{\\rm dec}$, which is a smoking-gun signature of dark radiation. It not only provides an evidence of dark radiation, but also sheds light on its production mechanism.  Some notes are in order.  We have assumed that the dark radiation anisotropic stress is induced only by the primordial GWs.  This is not in general true in the second order perturbation theory. Free-streaming particles (as well as other fluids) contribute to GWs at the second order in the scalar perturbation even if there is no primordial tensor perturbation.  However, this contribution is negligible for $r\\gtrsim 10^{-6}$~\\cite{Baumann:2007zm,Mangilli:2008bw}.  So far, we have considered dark radiation produced by the decay of $\\phi$.  However, it is possible that the dark radiation was once in thermal equilibrium and decoupled from thermal bath at the temperature $T_{\\rm dec}$.  In this case, the extra effective number of neutrino species is given by \\begin{equation} \\Delta N_{\\rm eff} = \\frac{4}{7}\\epsilon N_X \\left[ \\frac{10.75}{g_{*s}(T_{\\rm dec})} \\right]^{4/3}, \\end{equation} where \\begin{equation} \\epsilon = \\begin{cases} 1    & {\\rm ~~for~a~real~scalar},\\\\ 7/4 & {\\rm ~~for~a~chiral~fermion}, \\end{cases} \\end{equation} and $N_X$ counts the number of $X$ species.  If the decoupling temperature is higher than the weak scale, we need $N_X\\gtrsim 20$ for explaining $\\Delta N_{\\rm eff}\\simeq 1$.  The modulation in the GW spectrum, similar to the effect caused by of neutrinos apparent at the GW frequency of $10^{-10}$\\,Hz~\\cite{Weinberg:2003ur,Watanabe:2006qe}, appears at the frequency inside the range of DECIGO/BBO sensitivities for $T_{\\rm dec}\\sim 10^7$--$10^9$\\,GeV.  If the decoupling temperature is $\\mathcal O$(1)\\,MeV, $N_X\\sim 1$ is sufficient in order to obtain $\\Delta N_{\\rm eff}\\simeq 1$ but the dip in the GW spectrum cannot be seen in the GW detectors.  Instead, overall normalization of the GW spectrum at the observable frequency range, inferred from the measured tensor-to-scalar ratio, is enhanced by the factor $C_1\\sim 1.3$.  (At this epoch, dark radiation took part in thermal bath and there is no anisotropic stress damping on GW amplitudes with corresponding modes.)  This provides another indirect evidence of dark radiation."
    },
    "1208/1208.2979_arXiv.txt": {
        "abstract": "In cosmic ray air showers, the muon lateral separation from the center of the shower is a measure of the transverse momentum that the muon parent acquired in the cosmic ray interaction. IceCube has observed cosmic ray interactions that produce muons laterally separated by up to 400 m from the shower core, a factor of 6 larger distance than previous measurements.  These muons originate in high $p_T$ ($>$ 2 GeV/c) interactions from the incident cosmic ray, or high-energy secondary interactions. The separation distribution shows a transition to a power law at large values, indicating the presence of a hard $p_{T}$ component that can be described by perturbative quantum chromodynamics. However, the rates and the zenith angle distributions of these events are not well reproduced with the cosmic ray models tested here, even those that include charm interactions. This discrepancy may be explained by a larger fraction of kaons and charmed particles than is currently incorporated in the simulations. ",
        "introduction": "Introduction}  There have been many attempts to measure the cosmic ray composition at energies around and above the knee of the spectrum ($\\sim$ PeV) \\cite{gaisser90,Bluemer:2009zf}.  At these energies, direct measurements by balloon and satellite experiments have very limited statistics.  Ground based experiments rely on indirect measurements using observables such as the ratio of the measured electromagnetic energy to the number of muons {\\cite{IceCubecomposition,Antoni:2005wq}. These analyses are dependent on phenomenological calculations and simulations to relate the muon observations to an inferred composition; the result can be sensitive to the assumed hadronic interaction models \\cite{Kang:2010bs}.  Studies of high-energy ($\\gtrsim$ 1~TeV) muons with underground detectors have been an important part of this effort. The muons are produced early in air showers and probe the initial shower development \\cite{Klein:2009ew}. Two classes of muons are generally considered. ``Conventional'' muons come from pion and kaon decays, while ``prompt'' muons come from the decays of particles containing heavy quarks, mostly charm.  Conventional muons dominate at TeV energies, but, at energies above 100 TeV, prompt muons are expected to dominate \\cite{pasquali98}.  The resulting change in the slope of the muon energy spectrum has not yet been observed \\cite{Icecubenumu11}.  Studies of isolated muons, far from the shower core, can help understand the uncertainties due to phenomenological models.  Muon separations greater than about 30~m are largely due to the transverse momentum, $p_T$, imparted to the muon by its parent. For $p_T \\gtrsim 2$ GeV/c, these interactions can be described in the context of perturbative quantum chromodynamics (pQCD). Data from RHIC, the Tevatron, and the LHC are in quite good agreement with modern fixed order plus next-to-leading log calculations \\cite{vogt06}. These experimental studies give us some confidence in pQCD calculations for air showers.  Experimentally, the transition from soft interactions ({\\it i.e.} those with $p_{T} < 2$ GeV/c that are not describable in pQCD) to hard interactions is visible as a transition from a $p_T$ spectrum that falls off exponentially to one that follows a power law. At low $p_T$ the spectrum follows $\\exp{(-p_T/T)}$, with $T\\approx 220$ MeV/c for pions (somewhat higher for kaons and protons) \\cite{Adare:2011vy}. At higher $p_T$, the spectrum falls as $1/(1+p_T/p_0)^n$, where one fit found n=$13.0^{+1.0}_{-0.5}$ and $p_0 =1.9 ^{+0.2}_{-0.1}$ GeV/c \\cite{Adams:2004zg}. The transition is around 2 GeV/c for pions.  This spectral change should be visible in the lateral separation distribution.   The MACRO detector has previously measured the lateral separation between muons in air showers for primary energies ranging roughly from 10$^4$~GeV to 10$^6$~GeV \\cite{ambrosio99}. Buried under 3.8 km water equivalent of rock, MACRO has a minimum muon energy of about 1.3 TeV. MACRO measured muon pair separations out to a distance of about 65~m. Their simulations verified the linear relationship between $p_T$ and separation (with a small offset due to multiple scattering of the muons) out to a $p_{T}$ of 1.2~GeV/c, below the expected transition to the pQCD regime.  The muon $p_T$ is  related to the separation of the muon from the shower core by \\begin{equation} d_{T}\\eqsim\\frac{p_{T}Hc}{E_{\\mu}\\cos(\\theta)} \\label{eq:pt} \\end{equation} where $d_T$ is the perpendicular separation between the muon and the shower core, $H$ is the interaction height of the primary, $H/\\cos(\\theta)$ is the path length of the shower to the ground at a zenith angle $\\theta$, and $E_{\\mu}$ is the energy of the muon at generation. The interaction height of the parent of the muon is assumed to be synonymous with the primary interaction height, and the energy of the muon at generation is well approximated by its energy at the surface of the Earth.  In addition to the initial $p_T$, muons can separate from the shower core when they bend in the Earth's magnetic field or multiple scatter in the ice above the detector.  However, for the muon energies and separations considered here the gyroradius is on the order of 20,000 km and multiple scattering is similarly negligible; the initial $p_T$ is the dominant effect producing the separation.  A selection of muons with large transverse separation is biased toward events produced at high altitudes, high $p_T$, and low energy. The detector geometry imposes a minimum energy; most of the muons will naturally be near this threshold.  In the competition between altitude and $p_T$, the atmospheric density decreases exponentially with increasing altitude, while the $p_T$ should fall more slowly, only as a power law.  Although widely separated muons are biased toward high-altitude interactions, they still allow for studies of transverse momentum.  Events with large zenith angles are also preferred, since the muons are given more time to separate.  Of course, all of these factors should be appropriately modeled in the Monte Carlo.  In the IceCube neutrino telescope \\cite{Halzen:2010yj}, muons are detected with a 1 km$^3$ array of optical sensors buried in the Antarctic ice at depths between 1450 and 2450~m; individual muons studied in this analysis must have an energy at the Earth's surface of at least 400~GeV to reach the detector. The 125~m horizontal spacing between IceCube strings serves as a rough threshold for the minimum resolvable separation. For vertical muons with an interaction height of 50~km and an energy of 1~TeV, this corresponds to a $p_{T}$ of 2.5~GeV/c. The interaction height and muon energy vary from shower to shower, so the event-by-event uncertainty in $p_{T}$ approaches a factor of 2 if we assume average values for both.  The muon energy and cosmic ray interaction height for air showers that produce muons in the detector depend on the zenith angle. Figure \\ref{fig:intht} shows the interaction height as a function of true zenith angle for showers at sea level for DPMJET simulation (see Secs. \\ref{sec:sim} and \\ref{sec:disc} for a full description of the simulation). This dependence arises in part because the column depth of the atmosphere and the primary energy for air showers increases with zenith angle, leading to higher interaction heights for horizontal showers. The majority of showers have muons contained within 135~m of the shower core, but a small fraction have muons with larger lateral extensions; these showers interact much higher in the atmosphere. Figure \\ref{fig:minenergy} shows a fit to both the minimum and average energy of muons at the surface of the earth as a function of zenith angle, as well as the minimum energy calculated assuming continuous energy loss along the track ($dE/dx$). The energies shown are for the muon with the largest perpendicular distance from the shower core.   \\begin{figure} [htb] \\includegraphics[width=0.48\\textwidth]{paper_plot_separation_height.eps} \\caption{\\label{fig:intht} (Color online). The interaction height for all DPMJET simulated showers. Distributions for simulated showers with a true maximum muon separation less than 135~m and greater than 135~m are also shown.} \\end{figure}  \\begin{figure} [htb] \\includegraphics[width=0.48\\textwidth]{paper_plot_calc_emin_scatter.eps} \\caption{\\label{fig:minenergy} (Color online). The true energy at the surface of the Earth for the muon furthest from the shower core for simulated DPMJET shower events that pass all selection criteria versus zenith angle. Also shown are fits of the minimum and average energies. The values calculated using Eq. \\ref{eq:emin} are also shown.} \\end{figure}  The zenith angle has an impact on shower development. The 1450~m of ice above the detector shield it from vertical muons with energies less than about 400~GeV. For inclined showers, several effects come into play.  The distance between the target and detector rises, giving the muon more time to separate from the shower core.  However, the slant depth also increases, raising the muon energy threshold roughly exponentially with the ice thickness.  For average $dE/dx$ energy loss, the minimum energy at the Earth's surface is given by \\cite{chirkin04} \\begin{equation} E ({\\rm min}) = \\frac{a}{b}\\left[\\exp(Db/\\cos\\theta) - 1\\right] \\label{eq:emin} \\end{equation} where $D$ is the depth of the detector and $a$ and $b$ are constants that describe the energy loss of a muon in ice (with values of 0.177 GeV/mwe and 0.209 $\\times$ 10$^{-3}$/mwe, respectively) \\cite{chirkin04}.  The zenith angle distribution also depends on the parents of the muons.  At TeV energies, most muons originate from pions and kaons that decay before they interact. The probability of decay increases for larger zenith angles, because the pions and kaons spend more of their livetime at higher altitudes where the target density is lower so they are less likely to interact. The muon flux is \\cite{gaisser90,desiati10}  \\begin{align} \\frac{dN}{dE_\\mu} \\propto & \\frac{0.14 E_{\\mu}^{-2.7}}{\\rm cm^{2}\\,s\\,sr\\,GeV^{-3.7}} \\bigg[\\frac{1}{1+\\frac{1.1E_\\mu\\cos(\\theta)}{115 \\rm GeV}} \\nonumber\\\\ &+ \\frac{0.054}{1+\\frac{1.1E_\\mu\\cos(\\theta)} {850 \\rm GeV}} + \\frac{9.1\\times 10^{-6}}{1+\\frac{1.0E_\\mu\\cos(\\theta)} {5 \\times 10^{7} \\rm GeV}} \\bigg] \\label{eq:angulardist} \\end{align} where the three terms are for muons from pions, kaons, and charmed particles, respectively.  Charmed hadrons decay very quickly, leading to a flatter distribution in zenith angle.  This angular difference can be used to separate prompt muons from conventional muons; Fig. \\ref{fig:pidzen} shows the zenith angle of cosmic ray muons with energies of $\\sim$2~TeV produced by pion, kaon, or charm interactions. The fraction of muons from charm interactions increases for high zenith angles and higher energies. \\begin{figure} [htb] \\includegraphics[width=0.48\\textwidth]{paper_plot_cos_theta_ethresh_all.eps} \\caption{\\label{fig:pidzen} (Color online). The cosine of the zenith angle of DPMJET simulated muons produced by pion, kaon, or charmed particles interactions. The curves have been normalized to the peak bin and the muons energy is 2~TeV ($\\pm 10\\%$).} \\end{figure}   This paper extends the MACRO muon lateral separation measurements out to a separation of 400 m, well into the pQCD regime,  using 335 days of data collected with the partially completed IceCube detector.  The following sections give a description of the IceCube detector and an overview of the analysis.  The simulation is described in section \\ref{sec:sim} and the background reduction is discussed in section \\ref{sec:analysis}. The resulting distributions are discussed in sections \\ref{sec:results} and \\ref{sec:disc}. ",
        "conclusions": "IceCube has observed 34,754 muons with lateral separations between 135 m and 400 m; this corresponds to a transverse momentum of at least 2 GeV/c. The separation distribution is poorly fit by an exponential distribution with a ${\\chi}^2$/DOF of 61.5/21.  The fit improves when a power law component is included to a ${\\chi}^2$/DOF of 30.8/19, as expected from pQCD. However, the zenith angle distribution of the muons is unexpectedly flat, even when including the decay of charmed particles, and is poorly modeled by current simulations. This may be caused by an underproduction of kaons and charmed particles in the simulation. Future simulations with more sophisticated $p_{T}$ modeling may improve the disagreement in zenith angle.  IceCube has demonstrated the capability to resolve laterally separated muons in air showers. When improved simulation becomes available, future analyses could generate an estimate of muon parent ratios as well as a measurement of the transverse momentum spectrum in cosmic ray air showers."
    },
    "1208/1208.0099_arXiv.txt": {
        "abstract": "NEAT is an astrometric mission proposed to ESA with the objectives of detecting Earth-like exoplanets in the habitable zone of nearby solar-type stars. In NEAT, one fundamental aspect is the capability to measure stellar centroids at the precision of $5\\e{-6}$ pixel.  Current state-of-the-art methods for centroid estimation have reached a precision of about $4\\e{-5}$ pixel at Nyquist sampling. Simulations showed that a precision of 2 \u00b5-pixels can be reached, if intra and inter pixel quantum efficiency variations are calibrated and corrected for by a metrology system.  The European part of the NEAT consortium is designing and building a testbed in vacuum in order to achieve $5\\e{-6}$ pixel precision for the centroid estimation. The goal is to provide a proof of concept for the precision requirement of the NEAT spacecraft. In this paper we give the basic relations and trade-offs that come into play for the design of a centroid testbed and its metrology system. We detail the different conditions necessary to reach the targeted precision, present the characteristics of our current design and describe the present status of the demonstration. ",
        "introduction": "\\label{sec:INTRODUCTION}  \\subsection{Presentation of the NEAT concept}\\label{subsec:Presentation of the NEAT concept}  With the present state of exoplanet detection techniques, none of the rocky planets of the Solar System would be detected and indeed their presence is a very strong constraint on the scenarios of the formation of planetary systems. By measuring the reflex effect of planets on their central host stars, astrometry can lead us to the mass of planets and to their orbit determination. This technique is used frequently and is very successful to determine the masses and the orbits of binary stars. However it is necessary to go to space to reach the precision required to detect all planets down to the telluric regime. \\\\  We have been proposing a mission to ESA in the framework of the call for M missions in the Cosmic Vision plan which objective is to find most of the exoplanets of our Solar neighbourhood\\cite{Malbet11,Malbet12}. The objective is to use differential astrometry to complete the measurements obtained by other techniques in order to lower the threshold of detection and characterization down to the level of an Earth mass in the habitable zone of each system. We want to explore in a systematic manner all solar-type stars (FGK spectral type) up to 20 pc from the Sun. The satellite concept is based on formation flying technology with a satellite carrying a single primary mirror and another satellite carrying the focal plane (see Fig.~\\ref{neat_concept_diagram}). The measure is done using laser metrology and interferometry.  \\begin{figure}[t] \\begin{center} \\includegraphics[height = 40mm]{neat_concept_diagram_v3.pdf} \\caption{\\label{neat_concept_diagram}\\textbf{The proposed NEAT concept.} The metrology system projects dynamic Young fringes on the detector plane. The fringes allow a very precise calibration of the CCD in order to reach micro-pixel centroiding errors.} \\end{center} \\end{figure}  One of the fundamental aspects of the NEAT mission is the extremely high precision required to detect exo-Earths in habitable zone by astrometry. The amplitude of the astrometric signal that a planet leaves on its host star is given by the following formula: \\begin{equation}\\label{eq:astrometric_signal} A = 3 \\mu \\ml{as} \\times \\frac{M_{\\ml{Planet}}}{M_{\\ml{Earth}}} \\times \\left(\\frac{M_{\\ml{Star}}}{M_{\\ml{Sun}}}\\right)^{-1} \\times \\frac{R}{1\\ml{AU}} \\times \\left(\\frac{D}{1\\ml{pc}}\\right)^{-1} \\end{equation}  Where $D$ is the distance between the sun and the observed star, $M_{\\ml{Planet}}$ is the exoplanet mass, $R$ is the exoplanet semi major axis and $M_{\\ml{Star}}$ is the mass of the observed host star. For an Earth in the habitable zone located at 10 pc from the sun, the astrometric signal is 0.3 micro arcseconds (or $1.45\\e{-11}$ rad). This is smaller than the precision announced for the Gaia mission (launch scheduled for 2013) which should be 7 $\\mu$as, in optimal conditions. With a focal length of 40 meters, and taking into account a required signal to noise ratio\\cite{sim_double_blind_test09} of 6 and the required number of measurements per target \\cite{neat_number_of_measurements11}, the 0.3 $\\mu$as requirement to detect an Earth at 10 pc translates into a need to calibrate the pixelation error to $5\\e{-6}$ pixels for each integration, as shown by the NEAT error budget\\cite{neat_error_budget11}. \\\\  In the following subsections we will present the so-called centroid experiment, which goal is to demonstrate the feasibility of the latter requirement: first we present the theoretical aspects, then we introduce the JPL centroid experiment\\cite{Nemati11} and its CNES homologue.  \\subsection{Presentation of the centroid experiment: theoretical aspects}\\label{sec:centroid xp theoretical aspects}  Before entering further into the subject, we need to define what is measured, how we perform the measurements and how we define the errors we will be dealing with. We measure a distance between the location of several centroids. More precisely, we imitate the differential measurement technique of NEAT by comparing the position of the central target star to a reference frame defined by the position of the surrounding reference stars. The final error of the measure is comparable to the error of the centroid location estimation. A straightforward method to estimate a centroid location is to calculate the barycenter of the intensity of light. But this method is affected by the noise of the pixels located at large distances and prevent a precise estimation. \\\\  That is why the standard approach to precision centroid measurement is to do a least square fit of the PSF: a precise knowledge of the PSF is required. To mitigate this problem, we use a centroid displacement algorithm: one image is used to reconstruct the PSF, then the PSF is resampled at different locations. The location at which the resampled image best matches the second image is the estimation of the displacement between the two images. The key point here is that the PSF can be reconstructed very accurately because it is Nyquist sampled and because the pixels response functions are characterised by the metrology\\cite{Nemati11}. \\\\  The measure of the location of one centroid is affected by numerous types of errors:  \\begin{itemize} \\item \\textbf{Photon noise error:} Error due to the finite number of photon constituting the centroid. This is a fundamental limitation and represents the ideal error limit for a given centroid estimation method. A fundamental relation in our experiment is the precision reached for a given number of detected photo-electrons. If we knew the exact position of every photon detected during the integration (as if the pixels were all perfect and infinitely small and the detector had an infinite size), the best estimator of the centroid location would be the average location of all $N_{\\ml{ph}}$ photo-electrons, and would have a standard deviation of $\\frac{\\sigma_0}{\\sqrt{N_{\\ml{ph}}}}$ (by direct application of the Central Limit Theorem). In this relation $\\sigma_0$ is the deviation for one photon, i.e. it is the first moment of the centroid intensity distribution. This has important consequences in the design of our experiment. From this relation, we know that to reach a precision of $5\\e{-6}$ pixel, we have to use more than $N_{\\ml{ph}} = 7 \\e{9}$ photo-electrons. This number will determine the integration time necessary to reach the desired precision, given the flux on the CCD and the technical specifications of the latter. For practical purposes, we plan to be able to integrate in less than one hour: that will impose us strict design constraints. \\item\\textbf{Pixelation errors:} Errors related to the pixels. This term designate several types of errors. First, each pixel integrate the signal over a surface, and if the centroid is not properly sampled, it can introduce errors. Moreover, inter pixel variations also introduce some errors, even when the centroid is properly sampled: the pixels have different dark currents, sensitivities, they are not regularly spaced and the quantum efficiency (QE) profile within each pixel is not uniform. While the dark current and the sensitivity variations are usually calibrated with the dark and flat fields, at our level or precision we have to take into account all the parameters. \\item\\textbf{Truncation error:} Error caused by the finite size of the integration window: a part of the signal is inevitably lost and this leads to additional errors. \\end{itemize}  The final \\textbf{centroid error}, which is the standard deviation between the estimated and true position of the centroid, depends on all the errors mentioned above and on the centroid estimation algorithm used. If we assume that the errors are uncorrelated and the centroid estimation is optimal, we can express the centroid error as the root sum square of the other errors. This implies that all this sources of errors have to be kept below $5\\e{-6}$ pixel. Note that we have not considered position errors relative to the wavefront errors of the optics: as we do not change the relative positions of the mirror and the sources, they will remain identical between the measurements so we will be insensitive to them. \\\\  The method used to calibrate the CCD and to use the calibration information to perform the centroid measurements is described in Nemati et al.~(2011)\\cite{Nemati11} and Zhai et al.~(2010)\\cite{Zhai11}. Detailing the method is out of the scope of this article and we invite the reader to refer to the cited publications for further information.  \\subsection{JPL centroid experiment}\\label{sec:Presentation of the JPL centroid experiment}  An experiment aiming at demonstrating the feasibility of a state-of-the-art calibration at $5\\e{-6}$ pixels was conducted at the JPL (and it is still under development). The precision reached was $4\\e{\u22125}$ pixel for Nyquist sampled centroids. Figure~\\ref{fig:JPL_xp_picture} is a picture of the testbed and Fig.~\\ref{fig:JPL_xp_results} shows the result of the experiment.  \\begin{figure}[t] \\begin{center} \\subfigure[]{\\label{fig:JPL_xp_picture} \\includegraphics[width=8.2cm]{JPL_centroid_experiment.pdf}} \\hspace{5pt} \\subfigure[]{\\label{fig:JPL_xp_results} \\includegraphics[width=5.8cm]{JPL_results_differential.pdf}} \\end{center} \\caption{\\label{JPL_xp}\\textbf{JPL centroid experiment. Left: picture of the experiment. Right: result of the experiment.} The result displayed on the left is the Allan deviation of the measures of the distance between two centroids. One can see that the minimum deviation reached is $4\\e{-5}$ pixels for an integration time of a few minutes. Credit: Nemati et al.\\cite{Nemati11}} \\end{figure}  The CCD, the sources nor the mirror were moved during the whole integration time, so the dominant source of noise is initially the photon noise, until second order effects becomes dominant, after a few minutes of integration. The increasing values of the Allan deviation are caused by a drift of the centroids on the CCD. This drift is in turn thought to be the result of thermal fluctuations. As the calibration data from the metrology was not used in the centroid-fitting algorithm, a shift of the centroids has introduced errors related to inhomogeneities in pixel characteristics (mainly quantum efficiency and true locations). \\\\  During a future phase the CCD will be moved during the integration. By integrating over several sets of pixels the final calibration error should be averaged down. The results of the CCD calibration will be used to refine the fitting method as well.  \\subsection{CNES centroid experiment}\\label{subsec:Presentation of the CNES centroid experiment}  \\subsubsection{Context and timeline}  In order to strengthen the NEAT case in the next ESA call for M class missions, the European part of the NEAT consortium is designing and building a testbed very similar to the one that is used at the JPL. The main difference is that the metrology system will be made of integrated photonic components. The design of the testbed has begun on January 2012 and the components are now being procured. The bench will be assembled in September - October 2012. We expect to get some preliminary results at IPAG before 2013. Tests in vacuum should start in 2013. \\\\  The laboratories involved in this project are: the IPAG (Institut d'Astrophysique de Grenoble) - the laboratory where the experiment will initially take place, the CEA (Commisariat \u00e0 l'Energie Atomique et aux Energies Alternatives) where the electronics for the CCD camera are developed, the IAS (Institut d'Astrophysique Spatiale) where we have the possibility to use a vacuum chamber, the JPL (Jet Propulsion Laboratory) from which the past experience is very valuable to us. \\\\  The founding is done by the CNES (Centre National d'Etudes spatiales) and the labex OSUG@2020.  \\subsubsection{Presentation of the testbed}  The testbed is a simple optical bench that mimics the NEAT optical layout. A spheric mirror images five pinholes which are illuminated by a white source onto a CCD, so that the image is diffraction limited. The five pinholes represent stars, we will refer to them as \u2018\u2018pseudo stellar sources\". A set of single-mode fibres located at the edge of the mirror produce laser fringes on the detector. A schematic of the system's components is shown in Fig.~\\ref{fig:centroid_xp_diagram}. The optical set-up inside the vacuum chamber is shown in the Fig.~\\ref{fig:centroid_xp_optical_setup}.  \\begin{figure}[t] \\begin{center} \\subfigure[]{\\label{fig:centroid_xp_diagram} \\includegraphics[width=7cm]{centroid_xp_diagram.pdf}} \\hspace{5pt} \\subfigure[]{\\label{fig:centroid_xp_optical_setup} \\includegraphics[width=9cm]{optical_configuration.pdf}} \\end{center} \\caption{\\label{JPL_xp}\\textbf{Left: Schematic of the system's components. Right: The optical set-up of the experiment.}} \\end{figure}  The most innovative aspect of this experiment is the metrology system that will allow the micro-pixel calibration of the CCD. This system consist of at least two metrology bases (i.e the two pairs of single mode fibres), respectively aligned along the horizontal and vertical axis. The fibre extremities are located next to the mirror and project Young fringes on the detector. Additionally a phase modulator is used to dynamically sweep the fringes over the focal plane. By measuring the intensities variations of the signal for each pixel, one can characterise the inter and intra pixel response of the CCD and bring the centroid error down to the level of a few micro-pixels\\cite{Zhai11}. \\\\  In the next sections we present the specifications of the NEAT testbed, our design and finally we give the expected performances. For all these sections we will use the notations of the Table~\\ref{tab:notations}.  \\begin{table}[t] \\hspace{1cm} \\caption{\\label{tab:notations}Notations.} \\begin{center} \\begin{tabular}{|l|l|} \\hline Parameter & Notation\\\\ \\hline distance mirror to CCD (OA') & $L$\\\\ minimum/maximum wavelength of the pseudo stellar sources & $\\lambda_{\\ml{min}} / \\lambda_{\\ml{max}}$\\\\ wavelength used for the metrology & $\\lambda_m$\\\\ diameter of the entrance pupil & $D$\\\\ pixel size & $e$\\\\ mirror focal length (OF) & $f$\\\\ separation between the pseudo stellar sources (AB) & $s$\\\\ metrology baseline & $B$\\\\ metrology wavelength & $\\lambda_{\\ml{met}}$\\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We are in the process of building a testbed that will demonstrate the feasibility of measuring centroids to a precision of $5\\e{-6}$ pixel. This will strengthen the case for NEAT as it will show that astrometry down to sub-microarcsec precision is a valid technique for searching Earth-like exoplanets in the habitable zone of nearby stars.  The key specifications of a testbed capable of performing state-of-the-art centroid measurements are: Nyquist sampling of the airy spot, extreme stability of the optics and the use of an interferometric metrology system to calibrate the intra-pixel and inter-pixel quantum efficiency variations.  However, there are still a lot of developments to be made, the most critical ones being the theoretical developments related to the CCD calibration and the centroid estimation algorithms. We are currently developing an error budget that includes mechanical and thermal perturbations, the wavelength, the phase stability and the photon noise of the metrology.  We have presented here a optical bench that has been designed to demonstrate the capability to measure the distance between two sources at the level of 5e-6 pixels."
    },
    "1208/1208.0927_arXiv.txt": {
        "abstract": "{ It is a well established fact that some YSO jets (e.g. RW Aur) display different propagation speeds between their blue and red shifted parts, a feature possibly associated with the central engine or the environment in which the jet propagates. }{ In order to understand the origin of asymmetric YSO jet velocities, we investigate the efficiency of two candidate mechanisms, one based on the intrinsic properties of the system and one based on the role of the external medium. In particular, a parallel or anti-parallel configuration between the protostellar magnetosphere and the disk magnetic field is considered and the resulting dynamics are examined both in an ideal and a resistive magneto-hydrodynamical (MHD) regime. Moreover, we explore the effects of a potential difference in the pressure of the environment, as a consequence of the non-uniform density distribution of molecular clouds. }{ Ideal and resistive axisymmetric numerical simulations are carried out for a variety of models, all of which are based on a combination of two analytical solutions, a disk wind and a stellar outflow. The initial two-component jet is modified either by inverting the orientation of its inner magnetic field or imposing a constant surrounding pressure. The velocity profiles are studied assuming steady flows as well as when strong time variable ejection is incorporated. }{ Discrepancies between the speeds of the two oppositely directed outflows can indeed occur both due to unaligned magnetic fields and different outer pressures. In the former case, the asymmetry appears only on the dependence of the velocity on the cylindrical distance, but the implied observed value is significantly altered when the density distribution is also taken into account. On the other hand, a non-uniform medium collimates the two jets unevenly, directly affecting their propagation speed. A further interesting feature of the pressure-confined outflow simulations is the formation of static knots whose spacing seems to be associated with the ambient pressure. }{ Jet velocity asymmetries are anticipated both when multipolar magnetic moments are present in the star-disk system as well as when non-uniform environments are considered. The latter case is an external mechanism that can easily explain the large time scale of the phenomenon, whereas the former one naturally relates it to the YSO intrinsic properties. } ",
        "introduction": "\\label{sec:introduction}  Over the last few years, the two-component jet scenario emerges as a strong candidate for describing Young Stellar Object (YSO) outflows. Observational data of Classical T Tauri Stars (CTTS) (Edwards et al. \\cite{Edw06}; Kwan et al. \\cite{Kwa07}) indicate the presence of two genres of winds: one being ejected radially out of the central object (e.g. Sauty \\& Tsinganos \\cite{Sau94}; Trussoni et al. \\cite{Tru97}; Matt \\& Pudritz \\cite{Mat05}) and the other being launched at a constant angle with respect to the disk plane (e.g. Blandford \\& Payne \\cite{Bla82}; Tzeferacos et al. \\cite{Tze09}; Salmeron et al. \\cite{Sal11}). Consequently, CTTS outflows may be associated with either a stellar or disk origin, or with both outflow components having comparable contributions. In addition, such a scenario is supported by theoretical arguments (Ferreira et al. \\cite{Fer06}). An extended disk wind is required for the explanation of the observed YSO mass loss rates, whereas a pressure driven stellar outflow is expected to propagate in the central region, possibly accounting for the protostellar spin down (Matt \\& Pudritz \\cite{Mat08}; Sauty et al. \\cite{Sau11}). Numerical simulations have been also employed recently to investigate the various aspects of two-component jets (Meliani et al. \\cite{Mel06}; Fendt \\cite{Fen09}; Matsakos et al. \\cite{Tit09}, hereafter M09).  Detailed observations of some YSO with bipolar flows have shown peculiar velocity asymmetries between the blue and red shifted regions (e.g. Woitas et al. \\cite{Woi02}; Coffey et al. \\cite{Cof04}; Perrin et al. \\cite{Per07}). In particular, recent estimates of the RW Aur jet indicate that the speed of the approaching lobe is roughly 50\\% higher as compared to the receding one (Hartigan \\& Hillenbrand \\cite{Har09}). Although the above studies suggest that the asymmetry originates close to the source, observations of the same object carried out by Melnikov et al. (\\cite{Mel09}) point out a similar mass outflow rate for the two opposite jets, concluding that their different speeds are due to environmental effects. Thus, it is still an open question whether the discrepancy between the two hemispheres relies on an intrinsic property, such as the magnetic field configuration, or an external factor, such as the physical conditions of the surrounding medium.  Spectropolarimetric measurements of T Tauri stars suggest multipolar magnetospheres, with the dipolar component not always parallel to the rotation axis (Valenti \\& Johns-Krull \\cite{Val04}; Daou et al. \\cite{Dao06}; Yang et al. \\cite{Yan07}; Donati \\& Landstreet \\cite{Don09}). Non-equatorially symmetric field topologies may affect both the way that matter accretes on the protostar (Long et al. \\cite{Lon07}; Long et al. \\cite{Lon08}) and the jet launching mechanisms. For jets, it is not clear whether the magnetic field asymmetry persists for timescales comparable to the jet propagation timescales (decades-centuries). In fact, a periodic stellar activity could be directly associated with jet variability. In a similar context, Lovelace et al. (\\cite{Lov10}) simulated the central region of YSOs assuming complex magnetic fields. In one interesting case, where the magnetosphere is a combination of dipolar and quadrupolar moments, they find the extreme scenario of one-sided conical outflows being ejected from the star-disk interaction interface. On the other hand, the accretion disk could also have a quadrupolar magnetic field originating either from the dynamo mechanism or from advection (Aburihan et al. \\cite{Abu01}). In fact, it has been shown that non-bipolar disk field topologies can be a potential source for asymmetry in AGN jets (Wang et al. \\cite{Wan92}; Chagelishvili et al. \\cite{Cha96}).  The external medium may play a relevant role in the dynamical propagation of the outflow; we refer again to the detailed HST observations of RW Aur by Melnikov et al. (\\cite{Mel09}), whose findings can be consistent with the effects of inhomogeneities in the environmental conditions. Clumpy and filamentary structures are observed in molecular clouds over a wide range of scalelenghts, implying the presence of an external medium that surrounds the jets. In addition, an environment affecting jet propagation may be present but remain almost undetectable in observations, as discussed recently by Te\\c{s}ileanu et al. (\\cite{Tes12}). Theoretical arguments suggest that density anisotropies in the vicinity of YSOs could collimate YSO outflows and even induce oscillations in the jet's cross section (K\\\"onigl \\cite{Kon82}). Numerical simulations have investigated the effects of the surrounding gas, such as a collapsing environment whose ram pressure collimates a spherical wind (Delamarter et al. \\cite{Del00}) or an isothermal medium with a toroidal density distribution that results in a cylindrically shaped outflow (Frank \\& Noriega-Crespo \\cite{Fra94}; Frank \\& Mellema \\cite{Fra96}). Another class of simulations has studied the collimating role of a vertical outer magnetic field, whose magnetic pressure effectively confines an isotropic stellar wind (Matt \\& B\\\"ohm \\cite{Ma03a}; Matt et al. \\cite{Ma03b}). We note that although the external pressure is not thought to be responsible for the jet collimation observed at larger scales, it could still affect it and hence modify the propagation speed.  The goal of the present work is to study the feature of the velocity asymmetry within the context of the two-component jet models presented in M09. We take advantage of both theoretical and numerical approaches (Gracia et al. \\cite{Gra06}), setting as initial conditions a combination of two analytical YSO outflow solutions. Tabulated data of such MHD jets have been derived using the self-similarity assumption (Vlahakis \\& Tsinganos \\cite{Vla98}) and are available for both disk winds and stellar jets. We employ the Analytical Disk Outflow (ADO) solution from Vlahakis et al. (\\cite{Vla00}) and the Analytical Stellar Outflow (ASO) model from Sauty et al. (\\cite{Sau02}).  Matsakos et al. (\\cite{Tit08}) have addressed the topological stability, as well as several physical and numerical properties of each class of solutions separately. In M09 the two complementary outflows were properly mixed inside the computational box, such that the stellar outflow dominated the inner regions and the disk wind the outer. The stability and co-evolution of several dual component jet cases was investigated as a function of the mixing parameters and the enforced time variability. Furthermore, the introduction of flow fluctuations generated shocks, whose large scale structure demonstrated a strong resemblance to real YSO jet knots.  In this paper, two main scenarios are examined as possible mechanisms to produce velocity asymmetries. The first one refers to intrinsic YSO properties introducing distinct magnetic field topologies at each side of the equator. The second one refers to external effects arising from differences in the ambient pressure.  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{parallel_antiparallel.eps}} \\caption{Left panel: a quadrupolar disk field and a dipolar stellar magnetosphere. Right panel: a bipolar disk field and a quadrupolar stellar magnetosphere. The northern hemisphere of either case is referred to as a parallel configuration, whereas the southern as an anti-parallel. The dashed line indicates the surface that separates the disk wind from the stellar outflow. The simulations do not include the central region of the YSO, and hence we do not attempt to describe the star-disk interaction that takes place inside the grey ellipse. In addition, the requirement of $\\nabla\\cdot\\vec B = 0$ means that field lines must return along the accretion disk.} \\label{fig:parallel_antiparallel} \\end{figure*} In the first case, presented in Fig.~\\ref{fig:parallel_antiparallel}, each hemisphere is considered to have one of the following two magnetic field configurations: the field lines coming out of the stellar surface are parallel (north) or anti-parallel (south) to the large scale magnetic field. This implies a quadrupolar disk field and a dipolar magnetosphere (left), or equivalently a bipolar disk field surrounding a magnetic quadrupole (right).  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{external_pressure.eps}} \\caption{The northern hemisphere assumes a lower pressure, and hence a lower degree of collimation, as compared to the southern. The sketch is not in scale, the effects of this mechanism will hold at any distance wherein distinct environments are encountered. The magnetic field topology here is simple, a dipolar stellar magnetosphere is surrounded by a bipolar disk field.} \\label{fig:external_pressure} \\end{figure} In the second scenario, shown in Fig.~\\ref{fig:external_pressure}, the properties of the medium along the propagation axis are assumed to affect the collimation of the jet. Specifically, the outer pressure could modify the diameter of the outflow and in turn it would determine the wind speed. This mechanism does not depend on any YSO time scale, and it could appear at any height that the jet encounters heterogeneous environment. We note that although a thermal pressure of a hot tenuous medium is used in the simulations, this could equivalently represent the pressure due to a cold dense environment, turbulence or even a large scale magnetic field. Stute et al. (\\cite{Stu08}) truncated ADO solutions imposing a weaker outflow of the same type at outer radii. Here, we extend their study assuming a constant surrounding pressure instead, also with different values at each jet direction.  All our models focus on large enough scales avoiding the complicated dynamics of the star-disk interaction (e.g. Bessolaz et al. \\cite{Bes08}; Zanni et al. \\cite{Zan09}; Lovelace et al. \\cite{Lov10}). Essentially, these mechanisms are included merely as boundary conditions, an approach frequently adopted in the literature of jet propagation studies (e.g. Ouyed et al. \\cite{Ouy03}; Fendt \\cite{Fen06}; Te\\c{s}ileanu et al. \\cite{Tes09}).  Simulations are carried out to study the time evolution as well as the potential final steady states, comparing the results obtained between the two sides of the equatorial plane. We are especially looking for emerging asymmetries in the vertical velocity profiles. The resistive MHD regime is of particular interest (\\v{C}emelji\\'c et al. \\cite{Cem08}) since reconnection plays a key role, mainly at the locations where the magnetic field inverts sign. Furthermore, all cases are re-examined assuming that the flow has a strong time variable character, an effect that produces discontinuities along its axis. Note that we do not attempt to model the observed properties of the RW Aur jet, but rather to investigate in a general manner the intrinsic or extrinsic physical processes that could underlie the velocity asymmetry phenomenon.  The paper is structured as follows. Section \\S\\ref{sec:theory} describes the jet models and provides information on the numerical setup. Section \\S\\ref{sec:results} presents and discusses the obtained results while section \\S\\ref{sec:conclusions} reports the conclusions of this work. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we address the velocity asymmetries of YSO outflows by investigating two classes of candidate mechanisms that could possibly generate this feature.  The first class depends on the intrinsic properties of the YSO and assumes a parallel and an anti-parallel magnetic field configuration, one for each hemisphere. The application of physical reconnection results in different density and velocity distributions between the two sides, leading to observable speed discrepancies. This scenario can provide asymmetric jets coming from intrinsic physical conditions and processes, even without considering the complex dynamics of the star-disk interaction. The relative questions that arise are whether multipolar fields in the star-disk system can exist and survive for a time scale comparable to that of jet propagation, as well as what is the actual value of resistivity in YSO outflows, that could amplify or even suppress the asymmetry phenomenon.  The second class of mechanisms is based on external effects, namely when the YSO resides in a non-uniform environment. Imposing distinct outer pressures at some boundary line of the jet is found to directly affect the degree of collimation of each flow, which in turn results in significantly modified propagation speeds. This mechanism is appropriate to explain an outflow asymmetry on large time scales. In addition, static knots are found to manifest in the jet structure due to multiple recollimation locations. Their stability is robust despite the enforced flow perturbations, whereas their separation is found to be closely associated with the ambient pressure value. There is some evidence for stationary shocks in some systems (Matt \\& B\\\"ohm \\cite{Ma03a}; Bonito et al. \\cite{Bon11}; Schneider et al. \\cite{Sch11}), possibly associated with the collimation of the flow. However, shocks in protostellar jets are generally observed to propagate with the flow (Reipurth \\& Bally \\cite{Rei01} and references therein). The fact that standing recollimation shocks, predicted in non-variable simulations with a confining pressure, are not generally observed in YSO jets implies that either the dynamical evolution of the outflow is not strongly affected by the environment or that recollimation shocks are present but undetected.  In the case of the RW Aur jet, it is not clear enough from the available observations whether the asymmetry is associated with the central engine (Woitas et al. \\cite{Woi02}; Hartigan \\& Hillenbrand \\cite{Har09}) or with environmental effects (Melnikov et al. \\cite{Mel09}). Both scenarios could in principle be feasible. We have arbitrarily selected the values of our parameters to explore both candidate mechanisms in a general manner without attempting to model this particular object. In order to understand the applicability of a particular model to interpret the velocity asymmetry, more detailed observations are required, as well as more advanced numerical models, including 3D geometry and radiative processes."
    },
    "1208/1208.4115_arXiv.txt": {
        "abstract": "The observed properties of transiting exoplanets are an exceptionally rich source of information that allows us to understand and characterize their physical properties. Unfortunately, only a relatively small fraction of the known exoplanets discovered using the radial velocity technique are known to transit their host, due to the stringent orbital geometry requirements. For each target, the transit probability and predicted transit time can be calculated to great accuracy with refinement of the orbital parameters. However, the transit probability of short period and eccentric orbits can have a reasonable time dependence due to the effects of apsidal and nodal precession, thus altering their transit potential and predicted transit time. Here we investigate the magnitude of these precession effects on transit probabilities and apply this to the known radial velocity exoplanets. We assess the refinement of orbital parameters as a path to measuring these precessions and cyclic transit probabilities. ",
        "introduction": "The realization that we have crossed a technology threshold that allows transiting planets to be detected sparked a flurry of activity in this direction after the historic detection of HD~209458~b's transits \\citep{cha00,hen00}. This has resulted in an enormous expansion of exoplanetary science such that we can now explore the mass-radius relationship \\citep{bur07a,for07,sea07} and atmospheres \\citep{ago10,dem07,knu09a,knu09b} of planets outside of our Solar System. Most of the known transiting planets were discovered using the transit method, but some were later found to transit after first being detected using the radial velocity technique. Two notable examples are HD~17156~b \\citep{bar07} and HD~80606~b \\citep{lau09}, both of which are in particularly eccentric orbits. Other radial velocity planets are being followed up at predicted transit times \\citep{kan09a} by the Transit Ephemeris Refinement and Monitoring Survey (TERMS).  Planets in eccentric orbits are particularly interesting because of their enhanced transit probabilities \\citep{kan08,kan09b}. This orbital eccentricity also makes those planets prone to orbital precession. In celestial mechanics, there are several kinds of precession which can affect the orbital properties, spin rotation, and equatorial plane of a planet. These have been studied in detail in reference to known transiting planets, particularly in the context of the precession effects on transit times and duration \\citep{car10,dam11,hey07,jor08,mir02,pal08,rag09}. One consequence of these precession effects is that a planet that exhibits visible transits now may not do so at a different epoch and vice versa.  Here we present a study of some precession effects on known exoplanets. The aspect which sets this apart from previous studies is that we are primarily interested in planets not currently known to transit, particularly long-period eccentric planets which have enhanced transit probabilities and larger precession effects. We investigate the subsequent rate of change of the transit probability to show how they drift in and out of a transiting orientation. We calculate the timescales and rates of change for the precession and subsequent transit probabilities and discuss implications for the timescales on which radial velocity planets will enter into a transiting configuration, based upon assumptions regarding their orbital inclinations. We finally compare periastron argument uncertainties to the expected precession timescales and suggest orbital refinement as a means to measure this effect. ",
        "conclusions": "Transiting planets have become an essential component of exoplanetary science due to the exceptional opportunities they present for characterization of these planets. Many of the known exoplanets discovered through the radial velocity technique are currently not known to transit. However, transit probabilities can be substantially improved if the periastron argument approaches $\\omega = 90\\degr$. Since, for eccentric orbits, the periastron argument is time dependent as a result of their precession, planets which do not transit at the present epoch may transit in the future and vice versa. The planet Mercury falls quite central to the current distribution of calculated periastron precessions for the known exoplanets. This distribution has an eccentricity dependence but is most strongly affected by the orbital period. If a precession rate for a given planet is found to be markedly different from our calculations then this could be indicative of further, as yet undiscovered planets in that system. These additional planets would normally be detected from the radial velocity data unless insufficient observations allow them to remain hidden.  The periastron precession leads to a cyclic transit probability variation for all exoplanets with non-zero eccentricities. Timescales vary enormously but will likely lead to many of these planets transiting their host stars at some point in the future.  A reasonable question to ask at this point is if the periastron arguments of the known planets are known with sufficient precision to detect precession in any acceptable timeframe. Once again, we exploit the data extracted from the Exoplanet Data Explorer, described in Section \\ref{tranprob}. The uncertainties associated with the values of $\\omega$ for all these planets have a mean of $28\\degr$ and a median of $15\\degr$. This is much higher than the precession effects shown in Table \\ref{cyclictable}. A program of refining the orbits of the known exoplanets, such as that described by \\citet{kan09a}, would result in many of these precession effects to be detectable in reasonable time frames. For example the first planet in the table, HD~88133~b, has a precession rate that will cause a shift of $\\sim 0.3\\degr$ per decade. Uncertainties on $\\omega$ of less than one degree are not unsual and can certainly be achieved for those planets in particularly eccentric orbits. The exoplanet HD~156846~b has a current $\\omega$ uncertainty of $0.16\\degr$ \\citep{kan11} which demonstrates that such refinement is possible even for relatively long-period planets. More data and longer time baselines will produce subsequent improvements for many more planets which can result in the detection of the precession for high-precession cases.  The relevance of this work may be extended to the Kepler mission which has detected many candidate multi-planet systems \\citep{bor11a,bor11b,bat12}, most of which are likely to be real exoplanets \\citep{lis12}. Due to simply geometric transit probabilities, most of these systems will certainly have planets which are not transiting the host star at present. The known transiting multi-planet systems are largely in circular orbits, but may have periastron precession due to perturbations from other planets leading to an eventual transit from currently non-transiting planets in the system. For example, Kepler-19~c is known to exist from Transit Timing Variations of the inner planet, but does not currently have a detectable transit signature. Similarly, some of these planets will cease exhibiting an observable transit signature. Issues such as these are important for considering the completeness of these surveys in determining multi-planetary system architectures."
    },
    "1208/1208.6056_arXiv.txt": {
        "abstract": "We present the first blind interferometric detection and imaging of a millisecond radio transient with an observation of transient pulsar J0628+0909. We developed a special observing mode of the Karl G.\\ Jansky Very Large Array (VLA) to produce correlated data products (i.e., visibilities and images) on a time scale of 10 ms. Correlated data effectively produce thousands of beams on the sky that can localize sources anywhere over a wide field of view. We used this new observing mode to find and image pulses from the rotating radio transient (RRAT) J0628+0909, improving its localization by two orders of magnitude. Since the location of the RRAT was only approximately known when first observed, we searched for transients using a wide-field detection algorithm based on the bispectrum, an interferometric closure quantity. Over 16 minutes of observing, this algorithm detected one transient offset roughly 1\\arcmin\\ from its nominal location; this allowed us to image the RRAT to localize it with an accuracy of $1\\dasec6$. With \\emph{a priori} knowledge of the RRAT location, a traditional beamforming search of the same data found two, lower significance pulses. The refined RRAT position excludes all potential multiwavelength counterparts, limiting its optical luminosity to $L_{i'}<1.1\\times10^{31}$\\ erg s$^{-1}$ and excluding its association with a young, luminous neutron star. ",
        "introduction": "The fast ($\\lesssim1$\\ s) radio transient radio sky is populated by a wide range of source classes and physical processes. Neutron star pulses can be as brief as a nanosecond, making them the brightest objects in the universe and driving our understanding of coherent emission \\citep{2003Natur.422..141H}. Observations at fast time scales are sensitive to pulsars and magnetars \\citep{2000ApJ...541..367C}, extragalactic transients \\citep{2003ApJ...596..982M,2007Sci...318..777L}, stars \\citep{2008ApJ...674.1078O}, and planets \\citep{1999JGR...10414025F}. Propagation effects, like dispersion and scattering, are also detectable at GHz frequencies on millisecond time scales and can be used to probe the interstellar and intergalactic media \\citep{2002astro.ph..7156C}.  A new class of fast transient is the rotating radio transient \\citep[RRAT;][]{2006Natur.439..817M}. Like pulsars, RRATs are rotating neutron stars that emit radio pulses. Unlike pulsars, RRATs pulse rarely (roughly once out of every thousand rotations), so they are only detectable by their individual pulses. That observational distinction suggests that RRATs are physically distinct from ordinary pulsars \\citep[although that is not yet clear;][]{2011MNRAS.415.3065K}. Identifying optical and X-ray counterparts to radio-detected RRATs would allow us to study the parent neutron star and define their relation to ordinary pulsars \\citep{2009ApJ...702..692K}. Of the few dozen known RRATs \\citep{2009ApJ...703.2259D,2010MNRAS.401.1057K,2010MNRAS.402..855B}, only one has a clear association with an X-ray counterpart, suggestive of a magnetar-like object \\citep{2007ApJ...670.1307M,2009MNRAS.400.1439L}. However, since RRATs are an observationally defined class, it may be heterogeneous. Many more multiwavelength associations are needed before we can understand their underlying physics.  Large, single-dish telescopes have pioneered the study of RRATs and other fast radio transients through their high sensitivity and relatively simple signal processing requirements. Unfortunately, their design also limits their ability to survey efficiently, since survey speed scales as $N_{\\rm{a}}^2 \\, D^2$, where $N_{\\rm{a}}$\\ is the number of antennas used, and $D$\\ is their diameter \\citep{2008ASPC..395..225C,2011ApJ...742...12L}. More fundamentally, the arcminute-scale resolution of a large single-dish telescope is too coarse to find unique multiwavelength counterparts on arcsecond scales \\citep{2009MNRAS.400.1445K}. Pulsar timing can obtain very precise localizations, but requires months or years of monitoring. Single-dish telescopes are also susceptible to radio frequency interference (RFI), compromising some of their most exciting potential discoveries \\citep{2007Sci...318..777L,2010ApJ...715..939M}.  Interferometers, being composed of a distributed set of antennas, can overcome these limitations. They simultaneously have a large field of view and arcsecond resolution. Since interferometers are distributed over a large area, they are less susceptible to RFI and can localize interfering sources. However, pushing interferometric approaches down to subsecond time scales has seemed impractical due to the massive data throughput and computation needed to support image-based analysis. Novel signal processing concepts, such as the ``8gr8'' mode at the Westerbork Synthesis Radio Telescope \\citep{2009A&A...498..223J}, ``V-FASTR'' \\citep{2011ApJ...735...97W}, ``Fly's Eye'' \\citep{2012ApJ...744..109S}, or pulsar gating \\citep{2000ApJ...541..959B}, have opened access to more of the information available to interferometers. However, these concepts either require a priori information or are not based on visibilities, so they they lose some ability to blindly survey and localize transients.  Here we describe the first practical application of very high time-resolution observations with an interferometer, recording VLA visibilities every 10 ms, and using closure quantities to search for radio transients. Our goal was to improve the localization precision of RRAT J0628+0909 from arcminutes to arcseconds to facilitate a search for multiwavelength counterparts. Discovery observations with the Arecibo Observatory measured a pulse width of 10 ms, peak brightness of 85 mJy, and localization uncertainty of $3\\damin5$ \\citep{2006ApJ...637..446C,2009ApJ...703.2259D}, showing that it is accessible to our new VLA observing mode. Our observation marks the first use of the full VLA at this time scale and the first interferometric detection and localization of a fast radio transient. This also provides the first blind demonstration of a transient detection algorithm based on interferometric closure quantities \\citep{2012ApJ...749..143L} and a second example of how to implement a general-purpose, fast correlator \\citep{2011ApJ...742...12L}. We use the arcsecond localization of the transient to exclude its association with any cataloged multiwavelength counterpart, which excludes its association with a luminous, young, isolated neutron star. ",
        "conclusions": "We have demonstrated a new observing mode for the VLA that can write 10 millisecond visibilities, opening access to this exciting timescale to image and visibility analysis techniques. We used this new observing mode to find and image individual pulses from the poorly-localized RRAT J0628+0909. This is the first blind demonstration of our pulse detection algorithm based on the interferometric closure quantity known as the bispectrum and the first use of visibilities to detect and localize a fast radio transient.  We measure a new position for J0628+0909 of (6d28m36.25s, 9\\sdeg9\\arcmin14$\\dasec$9) (J2000) with an accuracy of roughly 1$\\dasec$6. Our interferometric localization is two orders of magnitude better than the localization made in its discovery observation. It is also consistent with new, sub-arcsecond localization made via a traditional timing analysis with Arecibo data spanning more than three years. Our localization for J0628+0909 excludes its association with any cataloged multiwavelength counterpart. SDSS imaging places an upper limit on the RRAT optical luminosity that excludes association with a young, luminous, isolated neutron star.  This paper presents the first real-world application of new signal processing and algorithms for using radio interferometers to the study millisecond transients. More than just pushing traditional slow transient techniques to fast time scales, our study of J0628+0909 shows how fast interferometric imaging encompasses and extends single-dish techniques. Future iterations can expand this concept by implementing real-time, all-time processing, opening the latent survey ability of radio interferometers."
    },
    "1208/1208.0736_arXiv.txt": {
        "abstract": "{Under the standard model extension (SME) framework, Lorentz invariance is tested in five binary pulsars: PSR J0737-3039, PSR B1534+12, PSR J1756-2251, PSR B1913+16 and PSR B2127+11C. By analyzing the advance of periastron, we obtain the constraints on a dimensionless combination of SME parameters that is sensitive to timing observations. The results imply no evidence for the break of Lorentz invariance at $10^{-10}$ level, one order of magnitude larger than previous estimation. ",
        "introduction": "\\label{sect:intro}  Unification of general relativity (GR) and quantum mechanics is a grand challenge in the fundamental physics. Some candidates of a self-consistent quantum theory of gravity emerge from tiny violations of Lorentz symmetry \\citep{Kostelecky2005,Mattingly2005}. To describe observable effects of the violations, effective field theories could be a theoretical framework for tests.  The standard model extension (SME) is one of those effective theories. It includes the Lagrange densities for GR and the standard model for particle physics and allows possible breaking of Lorentz symmetry \\citep{Bailey2006}. The SME parameters $\\bar{s}^{\\mu\\nu}$ control the leading signals of Lorentz violation in the gravitational experiments in the case of the pure-gravity sector of the minimal SME. By analyzing archival lunar laser ranging data, \\citet{Battat2007} constrain these dimensionless parameters at the range from $10^{-11}$ to $10^{-6}$, which means no evidence for Lorentz violation at the same level.  However, tighter constraints on $\\bar{s}^{\\mu\\nu}$ would be hard to obtain in the solar system because the gravitational field is weak there. Thus, for this purpose, binary pulsars provide a good opportunity. Because of their stronger gravitational fields, for example the relativistic periastron advance in the double pulsars could exceed the corresponding value for Mercury by a factor of $\\sim 10^5$, these systems are taken as an ideal and clean test-bed for testing GR, alternative relativistic theories of gravity and modified gravity, such as the works by \\citet{Bell1996}, \\citet{Damour1996}, \\citet{Kramer2006}, \\citet{Deng2009} and \\citet{Deng2011}.  Motivated by this advantage of binary pulsars, we will try to test Lorentz invariance under the SME framework with five binary pulsars: PSR J0737-3039, PSR B1534+12, PSR J1756-2251, PSR B1913+16 and PSR B2127+11C. In Sec. \\ref{sec:model}, the orbital dynamics of double pulsars in the SME will be briefed. Observational data will be used to constrain the SME parameters in Sec. \\ref{sec:obs}. The conclusions and discussions will be presented in Sec. \\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  In this work, we test Lorentz violation with five binary pulsars under the framework of standard model extension. It finds that $\\bar{s}_{\\omega}$, which is a dimensionless combination of SME parameters, is at the order of $10^{-10}$, whether all five systems are taken or top three systems with the smallest estimated uncertainties of periastron advances are used. This value, one order of magnitude greater than the estimation by \\citet{Bailey2006}, implies no evidence for the break of Lorentz invariance at $10^{-10}$ level.  Nevertheless, as mentioned by \\citet{Bailey2006}, the secular evolution of the eccentricity of the double pulsars should be included in the analysis. Its contribution is \\citep{Bailey2006} \\begin{equation} \\label{} \\bigg<\\frac{de}{dt}\\bigg> = \\frac{1}{e^3}n(1-e^2)^{1/2}(e^2-2\\varepsilon)\\bar{s}_e, \\end{equation} where \\begin{equation} \\label{} \\bar{s}_e = \\bar{s}_{PQ}-\\frac{\\delta m}{M}\\frac{2nae\\varepsilon}{e^2-2\\varepsilon}\\bar{s}_P. \\end{equation} $\\bar{s}_e$ is a combination of coefficients in $\\bar{s}^{\\mu\\nu}$ and sensitive to observations. However, there is lacking of timing observations on double pulsars running for a long enough time so that rare observations could show the secular change of $e$. Even though a few numbers could be derived from data, the uncertainties of them are quite larger than those of periastron advances.  Timing observations usually could set the upper bounds only, such as $|\\dot{e}|<1.9\\times10^{-14}$ s${}^{-1}$ for PSR B1913+16 \\citep{Taylor1989} and $|\\dot{e}|<3\\times10^{-15}$ s${}^{-1}$ for PSR B1534+12 \\citep{Staris2002}. Hence, we suppose that, at least in current stage, the constraints made by $\\dot{e}$ might be looser and the resulting upper bound is $|\\bar{s}_e|<3\\times 10^{-10}$. Although it is consistent with the values of $\\bar{s}_{\\omega}$ we obtain, the exact value of $\\bar{s}_e$ remains unknown. Therefore, unless timing observations could provide much more definitive results about $\\dot{e}$, the secular changes of eccentricity would not impose a tight constraint on $\\bar{s}^{\\mu\\nu}$ or combinations of $\\bar{s}^{\\mu\\nu}$.  Another issue for future work is to constrain the components of $\\bar{s}^{\\mu\\nu}$ directly with double pulsars. However, the choice of reference frame affects the values of these components so that a certain reference frame must be specified first and the projections of $\\bar{s}^{\\mu\\nu}$ will be along its standard unit basis vectors. For example, for comparing the constraints due to double pulsars and lunar laser ranging, $\\bar{s}^{\\mu\\nu}$ has to be projected along the same triad of vectors. It means the unit vectors $\\bm{P}$, $\\bm{Q}$ and $\\bm{k}$ (see Sec.\\ref{sec:model}) have to be decomposed in terms of these vectors, which requires the geometrical information of the orbit of the double pulsars, such as the orbital elements $\\Omega$ and $\\omega$. Unfortunately, timing observations are not sensitive to those two elements. This makes the components of $\\bar{s}^{\\mu\\nu}$ hard to access directly for now and demonstrates the advantages and availability of $\\bar{s}_{\\omega}$."
    },
    "1208/1208.5936_arXiv.txt": {
        "abstract": "We present high speed optical, spectroscopic and \\emph{Swift} X-ray observations made during the dwarf nova superoutburst of CC Scl in November 2011. An orbital period of 1.383 h and superhump period of 1.443 h were measured, but the principal new finding is that CC Scl is a previously unrecognised intermediate polar, with a white dwarf spin period of 389.49 s which is seen in both optical and \\emph{Swift} X-ray light curves only during the outburst. In this it closely resembles the old nova GK Per, but unlike the latter has one of the shortest orbital periods among intermediate polars. ",
        "introduction": "CC Scl was listed as RX J2315.5-3049 in the ROSAT X-ray catalogue (Voges et al.~1999), later identified with a 17.3 magnitude star, classified as a cataclysmic variable (CV) (Schwope et al.~2000), and found in the Edinburgh-Cape Survey where it became EC 23128-3105 (it is also in the 2MASS catalogue as 2MASS J23153185-3048476).  Its nature was confirmed when it was seen to rise in an outburst to magnitude 13.4 on 8 July 2000 (Stubbings 2000) followed by another about 100 days later. The latter was observed by Ishioka et al.~(2001) who suspected two apparent superhumps with a separation of 0.078 d. However, later photometric and spectroscopic observations made in quiescence (Augusteijn 2000; Chen et al.~2001; Tappert, Augusteijn \\& Maza 2004; Augusteijn et al.~2010) gave an orbital period $P_{orb}$ of 0.0584 d (1.402 $\\pm$0.005 h). Large amplitude flickering was noted and an occasional shallow narrow eclipse-like dip just after maximum light. Spectroscopic features included strong H and He\\,\\textsc{i} lines, with substantial strength of He\\,\\textsc{ii} 4686. Despite the last mentioned a classification simply as a dwarf nova was generally agreed upon.  Ishioka et al.~(2001) noted that at 9 d duration the apparent superoutburst of CC Scl was unusually brief:  $\\sim 14$ d is more to be expected of a short orbital period system. Following the two outbursts in 2000, CC Scl has had only seven outbursts reported by the AAVSO since (probably all rising to about the same maximum brightness), and only one outburst detected by the Catalina Real-Time Transient Survey (Drake et al.~2009) in the past 6 years -- the superoutburst of November 2011 discussed in this paper. Our observations were triggered by the CRTS announcement on 3 Nov 2011; PAW and BW were on the second night of a scheduled photometric observing run.  In Sect.~2 we describe our optical observations of CC Scl and analyse these to find persistent periodicities. In Sect.~3 we show the Swift X-ray satellite observations that were obtained and in Sect.~4 we discuss the results and compare them with other CVs, in particular the very similar behaviour of Nova Persei 1901 (GK Per). ",
        "conclusions": "\\subsection{CC Scl as an Intermediate Polar}  The list of definitely confirmed IPs (Mukai 2011) contains thirty-five members, to which we can now add CC Scl. Only five of those in Mukai's catalogue have $P_{orb} < 2$ h, i.e. are below the `orbital period gap'; CC Scl is thus an addition to this rare group. There are, however, two additional very short $P_{orb}$ systems that should be given full IP status: V455 And (Matsui et al. 2009; Silvestri et al. 2012) and WZ Sge (Warner \\& Pretorius 2008). In Tab.~\\ref{tableips} we list their properties, clustered into three groups: the slow rotators with $P_{orb}/P_{spin} \\sim 2$, intermediate with $P_{orb}$/$P_{spin} \\sim 10$ and those with $P_{orb}$/$P_{spin} \\sim 100$.  CC Scl is only the second example of the middle group (see Mukai (2011) for a complete plot of $P_{orb}$ versus $P_{spin}$ in which the rarity of low $P_{orb}$ systems is clearly visible). The first two groups represent respectively spins that are near equilibrium with angular momentum transferred directly from near to the inner Lagrangian point by strong magnetic fields, or, for weaker fields, after passage through an accretion disc which acquires angular momentum more appropriate to the circularisation radius (King \\& Wynn 1999). We deduce from this that the white dwarfs in HT Cam and CC Scl have magnetic moments significantly lower than those in the first four systems, and that V455 And and WZ Sge are lower still.  \\begin{table} \\centering \\caption{Intermediate Polars below the `orbital period gap'.} \\begin{tabular}{@{}lcccc@{}} Star & P$_{orb}$ (h) & $P_{spin}$ (s) & $P_{orb}$/$P_{spin}$ & $\\Delta T$ (d) \\\\[10pt] SDSS J2333+15   &      1.385   &         2500   &      1.99  &         -   \\\\ V1025 Cen         &      1.41    &         2147   &      2.36  &      discless \\\\ DW Cnc            &      1.44    &         2315   &      2.24  &        $2-4$  \\\\ EX Hya            &      1.64    &         4022   &      1.47  &        $2-3$  \\\\[5pt] CC Scl            &      1.38    &          389   &     12.8   &             9 \\\\ HT Cam            &      1.43    &          515   &     10.0   &             3 \\\\[5pt] V455 And         &  1.35        & 67.6         &  72   &    17 \\\\ WZ Sge           &  1.36        & 27.8         & 176   &    25 \\\\ \\end{tabular} \\label{tableips} \\end{table}    Tab.~\\ref{tableips} also notes dwarf nova-like outburst durations $\\Delta T$ of the group of short periods IPs (none reported in SDSS\\,J2333+15). V1025 Cen is believed to have a sufficiently strong field that no standard disc forms - instead mass transfer takes place through a centrifugally supported torus which does not suffer the standard disc instability. DW Cnc, EX Hya, and HT Cam have very brief infrequent outbursts, also unlike standard dwarf novae, which probably result from intermittent storage of gas just outside the magnetically defined inner edge of the accretion disc (Spruit \\& Taam 1993); therefore CC Scl is the only member of the middle group which has almost canonical disc instabilities, even though of slightly short durations for superoubursts. V455 And and WZ Sge have only superoutbursts.  The magnetosphere of the white dwarf in CC Scl will produce an inner disc radius of $r_0 \\sim 1.3 \\times 10^{10}$ cm, which is $\\sim 15 R(1)$ and $\\sim 0.6 r_d$, where the notation and equations 2.61, 2.83a, 7.17 of Warner (1995) have been used. Thus about 40\\% of the outer disc of CC Scl is available for regular dwarf nova outbursts (but only $\\sim 25$\\% of the outer disc of HT Cam). This may account for (a) the shorter than normal superoutbursts in CC Scl (there is less mass to drain out) and (b) the extremely short outbursts in HT Cam. These formulae do not apply to the longest $P_{spin}$ group in Tab.~\\ref{tableips} because the structure of what amounts to an accretion torus rather than a truncated disc is not known - but the effect evidently is to allow only very short-lived outbursts, and this may be what is happening in HT Cam as well.  From analysing X-ray spectra, Evans \\& Hellier (2005b) find that EX Hya, V1025 Cen and HT Cam are the only IPs that have low surface density accretion curtains in quiescence (with $N_H < 2 \\times 10^{21}$ cm$^{-2}$), implying a low rate of mass transfer in these systems, as is expected for CVs below the period gap. For CC Scl our X-ray spectra show $N_H > 4 \\times 10^{21}$ cm$^{-2}$ at spin minimum in outburst and quiescence, perhaps suggesting a higher accretion rate even though it too has an orbital period below the gap. Taking the quiescent luminosity of CC Scl from the latter part of Fig.~\\ref{ccsclunabs}, we have $L \\sim 3.6 \\times 10^{33} D^2$ erg s$^{-1}$, where $D$ is the distance in kpc. Considering the accretion rate to be $RL/GM$, and assuming a white dwarf mass of 0.6 M$_{\\odot}$, we find the accretion rate to be $\\sim 7 \\times 10^{-10} D^2$ M$_{\\odot}$ yr$^{-1}$. The high Galactic latitude of CC Scl ($b = -69^\\circ$) suggests  $D \\ll 1$, but without a better measure we cannot decide whether the accretion rate in CC Scl is high for its orbital period (for example EX Hya, with $P_{orb}$ close to that of CC Scl, has an accretion rate of $6 \\times 10^{-11}$ M$_{\\odot}$ yr$^{-1}$, Beuermann et al. 2003), and thus cannot test the absorption -- accretion rate relation suggested by Evans \\& Hellier.   \\subsection{Comparison with GK Per}  GK Per was Nova Persei 1901, one of the brightest novae of the twentieth century. It fell to its pre-eruption brightness after about 11 years, but since 1966 it has shown dwarf nova outbursts with an amplitude of 2 to 3 mag roughly every 3 years. In addition, it is established as an IP from a 351.3 s modulation in hard X-rays during outbursts (Watson, King \\& Osborne 1985), which is of low amplitude or absent during quiescence. However, it is seen at low amplitude in the photometric U band, even during quiescence, but is less coherent at longer wavelengths and better described as quasi-periodic near 380 s (Patterson 1991).  During outbursts the GK Per spectral lines stay in emission, rather than changing to broad absorption lines as in normal dwarf novae.  The similarities to CC Scl are marked, but one way in which GK Per differs totally from CC Scl is that it has an orbital period that is one of longest known for normal CVs: it is 1.997 d, deduced from spectroscopy.  This implies a large separation of components and a very large disc radius, of which only a small part of the central region is swept clear by the rotating magnetosphere of the primary. An effect of the large $P_{orb}$ is that it is difficult to make photometric observations that show any orbital brightness modulation, and even more difficult to see superhumps, if they exist. The possibility of superhumps arises because the secondary in GK Per is a subgiant reduced in mass through evolution and consequent mass transfer. The estimated mass ratio $q$ in GK Per is $0.55 \\pm 0.21$ (Morales-Rueda et al.~2002), so using the formulae $r_d/a = 0.6/(1+q)$ and $r_3 = 0.46a$ for the quiescent outer disc radius and the 3-to-1 resonance radii, respectively (Warner 1995), we have $r_3 \\sim 1.19 r_d$ for $q = 0.55$, and therefore the outer radius of the disc has only to expand by $\\sim$20\\% (to take up the angular momentum released by infalling matter) during a dwarf nova outburst for the disc to start evolving into an elliptical shape. The precession period $P_{prec}$ of the disc for $q = 0.55$ is $\\sim (10 - 12) P_{orb}$, i.e. $\\sim 20 - 24$ d in GK Per. Although there have been no obvious superoutbursts among the dwarf nova outbursts in GK Per, the 2006 outburst (Evans et al.~2009) was noted as unusually long ($\\sim 100$ d), and we note that there were three maxima spaced about 30 d apart, which might be evidence for an effect of $P_{prec}$. At $\\sim 60$ d the duration of most of the GK Per outbursts is too short for a $\\sim 30$-d periodicity to be properly established.  The X-ray, spectral and photometric similarities between GK Per and CC Scl suggest that the latter might have had a classical nova eruption sometime in the past few hundred years. There is no detectable nova shell in the available\u00e7 direct images, and GALEX images do not show any resolvable shell. An HST snapshot to look for H$\\alpha$ emission nebulosity could be justified.   \\subsection{Superoutbursts in intermediate polars}   Superoutbursts among IPs have been considered rare, perhaps even non-existent, e.g., Patterson et al.~(2011) have recently assessed the literature and conclude that although several IPs have dwarf nova outbursts, ``no confirmed superhumpers are known to be magnetic''. Nevertheless there are a couple of examples that might qualify: Patterson et al.~(2011) themselves cite TV Col and RX0704 as possible cases, but WZ Sge, the most extreme of the SU UMa stars, with outstanding superoutbursts and superhumps, has rapid oscillations ($\\sim 28$ s) that appear in optical and X-rays and which have been interpreted with an IP model (e.g. Warner, Livio \\& Tout 1996; Warner \\& Pretorius 2008). To this we can now certainly add CC Scl, as an unambiguous example of a superhumping IP, and from our earlier discussion GK Per could be a candidate. It is not clear from the long-term light curve of CC Scl (which, apart from the 2011 outburst, is very poorly sampled) whether all the outbursts except the latest are ordinary ones or whether they are superoutbursts. With two outbursts only 100 d apart it would be unusual to find they are both superoutbursts, but there is no sign in the light curve of a clear division in outburst ranges, and the fact that the only one (and of the usual maximum brightness), monitored with sufficient time resolution, turns out to have superhumps suggests that they all could be superoutbursts.  The question raised by CC Scl is why it alone, among the short orbital period IPs that have dwarf nova outbursts, is able to generate an elliptical disc and associated superhumps. The answer presumably lies in whatever limits outbursts in the other systems to only a couple of days in length - as a result there probably isn't sufficient time to get the eccentric mode fully excited."
    },
    "1208/1208.2985_arXiv.txt": {
        "abstract": "We investigate the impact of feedback -- from supernovae (SNe), active galactic nuclei (AGN) and a photo-ionizing background at high redshifts -- on the neutral atomic hydrogen (HI) mass function, the $b_{\\rm J}$ band luminosity function, and the spatial clustering of these galaxies at $z$=0. We use a version of the semi-analytical galaxy formation model {\\texttt{GALFORM}} that calculates self-consistently the amount of HI in a galaxy as a function of cosmic time and links its star formation rate to its mass of molecular hydrogen (H$_2$). We find that a systematic increase or decrease in the strength of SNe feedback leads to a systematic decrease or increase in the amplitudes of the luminosity and HI mass functions, but has little influence on their overall shapes. Varying the strength of AGN feedback influences only the numbers of the brightest or most HI massive galaxies, while the impact of varying the strength of photo-ionization feedback is restricted to changing the numbers of the faintest or least HI massive galaxies. {Our results suggest that the HI mass function is a more sensitive probe of the consequences of cosmological reionization for galaxy formation than the luminosity function. We find that increasing the strength of any of the modes of feedback acts to weaken the clustering strength of galaxies, regardless of their HI-richness. In contrast, weaker AGN feedback has little effect on the clustering strength whereas weaker SNe feedback increases the clustering strength of HI-poor galaxies more strongly than HI-rich galaxies. These results indicate that forthcoming HI surveys on next generation radio telescopes such as the Square Kilometre Array and its pathfinders will be exploited most fruitfully as part of multiwavelength survey campaigns.} ",
        "introduction": "Neutral hydrogen, both atomic (HI) and molecular (H$_2$), plays a fundamental role in galaxy formation, principally as the raw material from which stars are made. The amount of neutral hydrogen in a galaxy at any given time reflects the complex interplay between processes that either deplete it, such as supernovae, or replenish it, such as gas cooling from hot atmospheres surrounding galaxies and by mergers with other gas-rich galaxies. By quantifying the properties of HI in galaxies and noting how these properties vary with galaxy morphology and environment, we can glean insights into the physics of galaxy formation and test predictions of theoretical models.  Most of what we know about the HI properties of galaxies comes from surveys of the nearby Universe ($z\\lesssim 0.05$) such as HIPASS \\citep[HI Parkes All-Sky Survey; see][]{meyer.etal.2004} and ALFALFA \\citep[Arecibo Legacy Fast ALFA Survey; see][]{giovanelli.etal.2005}. Such surveys have revealed that HI-rich galaxies tend to be late-type \\citep[e.g.][]{kilborn.etal.2002,evoli.etal.2011}, but it is common for early-type field galaxies to host HI \\citep[e.g.][]{serra.etal.2012}; that the HI mass function is well described by a Schechter function \\citep[e.g.][]{zwaan.etal.2003,martin.etal.2010} but its shape depends on environment \\citep[e.g.][]{zwaan.etal.2005,springob.etal.2005,kilborn.etal.2009}; and that HI-rich galaxies are among the most weakly clustered galaxies known \\citep[e.g.][]{meyer.etal.2007,basilakos.etal.2007,passmoor.etal.2011}. However, knowledge of the HI properties of galaxies will improve many-fold over the coming decade with the advent of next generation HI galaxy surveys on ASKAP \\citep[cf.][]{askap.science.2008}, MeerKAT \\citep[cf.][]{meerkat.2007}) and ultimately the SKA itself \\citep[e.g.][]{baugh.etal.2004,blake.etal.2004, power.etal.2010,abdalla.etal.2010,kim.etal.2011}.  Future galaxy surveys on the SKA and its pathfinders are expected to revolutionize our view of the HI Universe and so it is timely to ask precisely what these surveys can teach us about the physical processes that drive galaxy formation. In this paper we focus on feedback -- from stars in the form of winds driven by supernovae (SNe), from accreting super-massive black holes in the form of active galactic nuclei (AGN) heating, and from a photo-ionizing background in the early Universe -- and evaluate how the global properties of HI in galaxies such as the HI mass function, which quantifies the number density of galaxies of a given HI mass, and the 2-point correlation function and halo occupation distributions \\citep[HODs; e.g.][]{kim.etal.2011}, which quantify spatial clustering, are shaped by different sources of feedback.  In galaxy formation models, the strengths of these processes and the interplay between them have been traditionally set by examining the predictions of the galaxy luminosity function in the $b_{\\rm J}$ band, so we will also study how this statistic changes. To do this, we use the version of the semi-analytical galaxy formation model {\\texttt{GALFORM}} of \\citet{cole.etal.2000} as it has been developed by \\citet{lagos.etal.2011a}; this calculates self-consistently the HI properties of galaxies by splitting their interstellar media (ISM) into HI and H$_2$ phases using the empirical relations of \\citet{blitz.2006} and \\citet{leroy.etal.2008}, and links star formation in galaxies not to their cold gas masses, as assumed in previous models \\citep[cf.][]{cole.etal.2000,baugh.2006}, but to their H$_2$ masses, as suggested by recent observations \\citep[e.g.][]{bigiel.etal.2008}. As shown in \\citet{lagos.etal.2011b}, this model reproduces the observed HI mass function at $z$=0, accurately reproducing its amplitude and shape at intermediate and low HI masses.  This is particularly interesting because we expect feedback to influence the global properties of galaxies selected either by their stellar mass or light. Previous studies have shown that SNe are pivotal in fixing the amplitude and slope of the luminosity function \\citep[e.g.][]{cole.etal.2000,benson.etal.2003a} while AGN heating suppresses the formation of massive galaxies and dictates the form of the bright end \\citep[cf.][]{bower.etal.2006,croton.etal.2006}. In contrast, the influence of SNe and AGN on the HIMF appears more subtle. For example, \\citet{power.etal.2010} found that galaxy formation models that included or excluded AGN heating (such as \\citealt{delucia.blaizot.2007} and \\citealt{baugh.etal.2005} respectively) predict HI mass functions that reproduce equally well the observed $z\\!\\simeq\\!0$ data. The HI mass function offers the possibility of placing stronger constraints on the strength of feedback in semi-analytical models, which previously used only the optical properties of galaxies \\citep[e.g.][]{benson.etal.2003a,bower.etal.2010}. At the same time, the variation of clustering strength with galaxy properties provides us with important clues about the physics of galaxy formation. Any discrepancy between observational measurements of clustering and the predictions of galaxy formation models indicates the need to improve the models, either by refining the modelling of the physical processes included or by considering the addition of further effects \\citep[see, for example,][]{kim.etal.2009}.  The structure of the paper is as follows. In \\S\\ref{sec:model}, we provide a brief overview of the \\citet{lagos.etal.2011a} model and describe how the different modes of feedback (SNe, AGN and photo-ionization) are implemented in {\\texttt{GALFORM}}. In \\S\\ref{sec:himf} we show how the different forms of feedback influence both the galaxy luminosity function in $b_{\\rm J}$ band and the HI mass function, {and we deduce using likelihood maximization precisely what information we glean from them.} In \\S\\ref{sec:clustering} we investigate how the spatial clustering of HI-poor and HI-rich galaxies are influenced by feedback by inspecting the 2-point correlation function and HODs. Finally, in \\S\\ref{sec:summary}, we summarize our results and discuss their implications for testing the modelling of feedback with forthcoming HI surveys. ",
        "conclusions": "\\label{sec:summary}  We have used the version of the {\\texttt{GALFORM}} semi-analytical galaxy formation model \\citep[cf.][]{cole.etal.2000} that has been extended by \\citet{lagos.etal.2011a} to study how feedback from SNe, AGN and photo-ionization shapes both the galaxy luminosity and HI mass functions, and the spatial clustering of HI galaxies. The advantage of the \\citet{lagos.etal.2011a} model is that it is one of the first to predict explicitly and in a self-consistent manner how much HI should reside in a galaxy at a given time \\citep[see also][]{fu.etal.2010} and it has reproduced successfully the observed HI mass function at low-to-intermediate masses \\citep[cf.][]{lagos.etal.2011b}. The main results of our study can be summarised as follows;  \\vspace{0.4cm}  \\noindent$\\bullet$ Feedback from SNe regulates the amplitude and shape of both the galaxy luminosity function and the HI mass function. The effect is systematic, with more (less) efficient SNe feedback leading to fewer (more) galaxies of a given luminosity or HI mass. It also regulates the amplitude of the 2-point correlation function of HI-rich galaxies at small separations and the numbers of galaxies per halo mass as quantified by the HOD, regardless of their HI richness.  \\vspace{0.2cm}  \\noindent$\\bullet$ Feedback from AGN has little effect on either the faint-to-intermediate luminosity end of the galaxy luminosity function or low-to-intermediate mass end of the HI mass function -- its impact is strongest for the brightest galaxies, as previous studies have argued \\citep[e.g.][]{croton.etal.2006,bower.etal.2006}, but it also affects the number of highest HI mass galaxies. A similar result has been reported by \\citet{fabello.etal.2011}, and it also helps to explain the cold gas mass-halo mass relation presented in Fig. 3 of \\citet{kim.etal.2011} for the \\citet{bower.etal.2006} model. In contrast, its impact on the clustering of HI galaxies is minor or negligible.  \\vspace{0.2cm}  \\noindent$\\bullet$ Feedback from photo-ionization is most pronounced for the faintest and most HI-poor galaxies. The redshift at which reionization occurs is not important -- but the mass scale, as measured by $V_{\\rm cut}$, is, and its influence is evident in both the shape and amplitude of the faint end of luminosity function and the low-mass end of the HI mass function. Our analysis suggests that the low-to-intermediate mass end of HI mass function offers the potential to constrain the models of photo-ionization. Interestingly, we find that the locally measured HI mass function can constrain the minimum mass of dark matter halos that could have hosted galaxies that contributed to reionization. This complements existing studies of the high redshift Universe that have used, for example, the electron scattering optical depth \\citep[e.g.][]{choudhury.etal.2008},  the Lyman-$\\alpha$ forest and \\citep[e.g][]{srbinovsky.wyithe.2010} and the UV luminosity function \\citep[e.g.][]{munoz.loeb.2011} to estimate the minimum halo mass. Our analysis suggests that the circular velocities of halos in which galaxy formation was suppressed cannot be larger than $\\sim\\!70\\,\\rm km/s$. \\vspace{0.2cm}  \\noindent$\\bullet$ {Strong modes of feedback, as they are implemented in {\\texttt{GALFORM}}, act to suppress the clustering strength of galaxies, regardless of their HI-richness. HI-poor galaxies cluster more strongly than HI-rich galaxies if the strength of SNe feedback is weakened, whereas there is little change in the clustering of either HI-rich or HI-poor galaxies if the strength of AGN feedback is weakened.}  \\vspace{0.4cm}  \\noindent Our study suggests that forthcoming HI galaxy surveys can be most fruitfully exploited scientifically if they are part of a multiwavelength campaign. Such campaigns are being planned; for example, DINGO (Deep Investigation of Neutral Gas Origins) on ASKAP \\citep[cf.][]{meyer.2009} will probe the galaxy population out to $z \\lesssim 0.4$ and will be combined the GAMA (Galaxy And Mass Assembly) survey \\citep[cf.][]{driver.etal.2011}. These data-sets will allow us to study not only the HI properties of galaxies that host AGN or young starbursts, but also the effects of environment and evolution with redshift."
    },
    "1208/1208.0452_arXiv.txt": {
        "abstract": "{LHS\\,1070 is a nearby multiple system of low mass stars. It is an important source of information for probing the low mass end of the main sequence, down to the hydrogen-burning limit. The primary of the system is a mid-M dwarf and two components are late-M to early L dwarfs, at the star-brown dwarf transition. Hence LHS\\,1070 is a valuable object to understand the onset of dust formation in cool stellar atmospheres.} {This work aims at determining the fundamental stellar parameters of LHS\\,1070 and to test recent model atmospheres: \\mbox{BT-Dusty}, \\mbox{BT-Settl}, DRIFT, and MARCS models.} {Unlike in previous studies, we have performed a $\\chi{^2}$-minimization comparing well calibrated optical and infrared spectra with recent cool star synthetic spectra leading to the determination of the physical stellar parameters $T_\\mathrm{eff}$, radius, and $\\mathrm{log}\\,g$ for each of the three components of LHS\\,1070. } {With exception of the MARCS models which do not include dust formation, the models are able to reproduce the observations and describe the main features of the visible to IR spectra. This is consistent with the fact that dust formation prevails in the B and C component atmospheres.  The parameters obtained with the DRIFT models confirm the values determined in earlier studies.   But important differences between models are observed, where the MARCS model is too bright in the $H$ and $K$ bands, and the BT-Settl and BT-Dusty models systematically yield up to 100\\,K higher $T_\\mathrm{eff}$ in the case of the B and C components. This confirms a trend for models without, or with less efficient cloud formation, to predict higher $T_\\mathrm{eff}$ than models richer in dust (DRIFT). Even models including cloud physics however still produce slightly too bright $J$ band flux, showing as too blue $J-K$ colors. The onset of dust formation remains therefore a particularly challenging regime to understand.} {} ",
        "introduction": "The lower end of the Hertzsprung-Russel diagram has much importance as the vast majority of stars in the Galaxy are low mass stars. In the Galaxy, 70\\% of the stars are M dwarfs. They contribute over 40\\% of the total stellar mass content \\citep{Gould1996,Mera1996,Henry1998}. These M dwarfs have a mass that ranges from $0.6 M_\\odot$ to the hydrogen burning limit of about 0.075 to $0.085 M_\\odot$ depending on the metallicity \\citep{Chabrier2000b}. These stars are found in any population, from young metal rich M-dwarfs in open clusters \\citep{Reid1993,Leggett1994} to the several billion years old metal poor dwarfs in the galactic halo (Green and Morgan 1994) and in the globular clusters \\citep{Cool1996,Renzini1996}. Such low mass stars are an important probe for our Galaxy as they carry fundamental information regarding the stellar physics, galactic structure and formation, and its dynamics. In addition, the existence of brown dwarfs or planets being discovered and confirmed around M-dwarfs \\citep[][and references therein]{Butler2004,Bonfils2011} plays an important role in understanding the formation of brown dwarfs and planets.  Despite their large number in the Galaxy, little is known about low mass stars because of the difficulty (i) to get a homogeneous sample with respect to the age and metallicity due to their intrinsic faintness, and (ii) to disentangle the parameter space ($T_\\mathrm{eff}$, $\\mathrm{log}\\,g$, and metallicity). Indeed a number of studies have shown that a change in temperature or gravity can compensate for a change in metallicity to some degree. An additional difficulty is the complexity of their atmospheres: convection in optically thin regimes, molecules, and dust cloud formation for the later types. Water vapor and CO bands dominate the Rayleigh-Jeans branch of the spectral energy distribution at infrared wavelengths ($> 1.3\\, \\mu$m), while TiO, VO, and metal hydrides govern the corresponding visual ($> 4000\\AA$) to near-infrared ($< 1.3\\, \\mu$m) spectral energy distribution.  Convection reaches out to the optically thin (as far as to an optical depth $\\tau \\sim 10^{-3}$) portion of the atmospheres, flattening the temperature gradient of the atmosphere \\citep[]{Allard1997}. \\cite{Ludwig2002,Ludwig2006} have determined the mixing length based on a comparison of the mixing length theory used in 1D static model atmospheres and Radiation Hydrodynamic simulations. In M dwarfs later than M6 the outermost temperatures fall below the condensation temperatures of silicate grains, which leads to the formation of dust clouds \\citep[see e.g.][]{Tsuji1996a,Tsuji1996b, Allard1997,Ruiz1997,Allard1998a,Allard1998b}. These processes complicate the understanding of these cool atmospheres. \\\\ One approach to study the physics at the low end of the main sequence is to compare observed spectra with synthetic spectra from various authors and modelling techniques. The determination of the physical parameters (effective temperature, gravity, metallicity, radius) is obtained by spectral synthesis, i.e. $\\chi^2$ minimization.  LHS\\,1070 is a low mass multiple system of cool dwarfs discovered by \\cite{Leinert1994}, with visual magnitude 15. It is located at a distance of 7.72$\\pm$0.15 pc from the Sun \\citep{Costa2005} and is considered as a member of the disk population with a probable age of around 1 Gyr \\citep{Reiners2007a}. The spectral types for the A, B, and C components were found to be M5.5-M6, M8.5, and M9-M9.5, respectively \\citep{leinert2000}. A fourth component was suspected very close to the primary by \\cite{Henry1999} from HST Fine Guidance observations, but this detection is no longer considered to be real (T. Henry, private communication). The latest orbit determination has been performed by \\cite{K2012}, with semi major axes of 0.458'' for the close pair BC and 1.112'' for the wide orbit of BC around component A. E.g., on December 12, 2003, component B was separated from A by 1.77\" at position 13$^{\\circ}$, and component C from B by 0.41\" at 178$^{\\circ}$.  \\citet{Leinert1998,leinert2000} have derived effective temperatures of 2950\\,K, 2400\\,K and 2300\\,K for the components based on spectral analysis, and found that B and C showed clear signatures of dust in their spectra. They presented photometric mass estimates ranging from 0.109 to $0.079 M_\\odot$ for the three stars based on theoretical isochrones, thus reaching right down to the minimum hydrogen burning mass. This mass range makes LHS\\,1070 a valuable system for understanding the formation of dust in cool atmospheres and the processes that occur at the star/brown dwarf transition. LHS\\,1070 is therefore a testbed to validate and define further developments of both atmospheric and interior models at  the lower end of the main sequence.  We assume the same age and composition for the three components of this system for simplicity.  In this paper, we present the spectral synthesis of the components A, B, and C. We determine their physical parameters by comparing the well-calibrated HST spectra in the optical (from FOS) and in near and mid-IR (from HST/NICMOS and ISOPHOT-S) with synthetic spectra computed from recent stellar atmosphere models: BT-Dusty and BT-Settl \\citep[]{Allard2012a}, MARCS \\citep{Gustafsson2008}, and DRIFT \\citep{Witte2009}. Observations and data reduction are described in Sec.~\\ref{obs}. Sec.~\\ref{mod} presents the atmosphere models used in the analysis. In Sec.~\\ref{result} we give the determination of the stellar parameters and show the comparison between observed and modeled spectra. Discussion and conclusion follow  in Sect.~\\ref{ccl}. ",
        "conclusions": "\\label{ccl}  This paper presents the results from spectral synthesis analysis for the LHS\\,1070 triple system. This system has been extensively observed from the optical to the IR, and dynamical masses have been determined \\citep{Leinert2001,Seifahrt2008}. Therefore, it constitutes a testbed of model atmospheres of low-mass stars. Band strength indices are used to measure TiO, CaH and PC3 features to classify their spectral type. The components are classified as M5.5, M9.5, and L0 dwarfs, and their atmospheres lie in a temperature range where dust starts to form. We have determined the physical parameters $T_\\mathrm{eff}$, $\\mathrm{log}\\,g$, metallicity, and radius for  the three components of the LHS\\,1070 system by comparing the observed spectra with the synthetic spectra computed with the most recent atmospheric models: BT-Dusty, BT-Settl, MARCS, and DRIFT. All the models agree for a solar metallicity for the system. The derived gravity is 5.0 dex and agrees within the uncertainties with the values derived from the dynamical mass \\citep{Seifahrt2008}.  We found the same value for $T_\\mathrm{eff}$ for the primary from all models while differences of 100\\,K and 200\\,K are found for components B and C depending on the dust density content of the model atmosphere used. The revised oxygen abundance by \\cite{Asplund2009} and \\cite{Caffau2011} yield significant improvements of the BT-Settl fits to the primary compared to earlier studies based on the larger values of the solar oxygen abundance of \\citet{Grevesse1998}. These improvements are described in \\cite{Allard2012a}. The even lower abundances of \\citet{Grevesse2007} used in the MARCS models lead to an excess of near-IR flux due to weaker water vapor absorption.  The DRIFT and BT-Settl models differ mainly in their numerical approach in solving the equations for grain growth, sedimentation and opacities: the DRIFT model solves them from the top to the bottom of the atmosphere, while the BT-Settl model solves them from the bottom to the top of the atmosphere. This causes the BT-Settl model to tend to have a deficit of grains in the upper atmospheric layers compared to the DRIFT model despite an adequate account in both models of supersaturation effects. Despite these fundamental differences, the resulting grain sizes obtained by the models are quite similar.  The over-luminosity shared by the models in the $J$ bandpass could be indicative of grains of larger sizes and/or more numerous in the LHS1070 B and C component  atmospheres. The results confirm the \\cite{Allard2012a} findings based on $T_\\mathrm{eff}$-color constraints.  The $T_\\mathrm{eff}$ values found with the DRIFT atmospheres for components B and C agree with the \\cite{leinert2000} findings. Except for the MARCS models which do not include dust treatment, the models are able to reproduce the observations and describe the main features of the visible to IR spectra for all three components. This raises the confidence level in the dust-modelling approach.  However, the calculation of opacities for composite grains relies on relatively simple approximations and also does not account for possible distributions of grain shapes and structures such as porosity. Both models rely on results of radiation hydrodynamical simulations that provide the mixing and overshooting which compensates sedimentation effects. One problem is certainly that the translation of the resulting radial velocity field into a diffusion coefficient is currently uncertain.  It is also possible that the mixing effects are being currently underestimated by local 2D simulations, and that additional mixing is provided by other phenomena on larger scales such as global rotation effects."
    },
    "1208/1208.1986_arXiv.txt": {
        "abstract": "We present a study exploring a systematic effect on the brightness of type Ia supernovae using numerical models that assume the single-degenerate paradigm. Our investigation  varied the central density of the progenitor white dwarf at flame ignition, and considered its impact on the explosion yield, particularly the production and distribution of radioactive \\Ni{56}, which powers the light curve.  We performed a suite of two-dimensional simulations with randomized initial conditions, allowing us to characterize the statistical trends that we present. The simulations indicate that production of Fe-group material is statistically independent of progenitor central density, but the mass of stable Fe-group isotopes is tightly correlated with central density, with a decrease in the production of \\Ni{56} at higher central densities. These results imply progenitors with higher central densities produce dimmer events. We provide details of the post-explosion distribution of \\Ni{56} in the models, including the lack of a consistent centrally-located deficit of \\Ni{56}, which may be compared to observed remnants. By performing a self-consistent extrapolation of our model yields and considering the main-sequence lifetime of the progenitor star and the elapsed time between the formation of the white dwarf and the onset of accretion, we develop a brightness-age relation that improves our prediction of the expected trend for single degenerates and we compare this relation with observations. ",
        "introduction": "\\label{sec:intro}  Type Ia supernovae (SNeIa; singular SNIa) are bright, transient astronomical events identified by a peak-light spectrum showing no evidence of hydrogen but absorption lines of singly-ionized silicon~\\citep{Mink41, Fili97}.  These events follow from explosive thermonuclear burning of degenerate stellar material composed principally of C and O, which synthesizes $\\sim\\!\\!0.6 \\Msol$ of radioactive \\Ni{56}.  The decay of this \\Ni{56} powers the light curve~\\citep{truranetal67, colgatemckee69, arnett:type, pinto.eastman:physics}.  The progenitor systems of these explosions remain the subject of considerable debate and active research.  Observations, however, indicate these events largely form a homogeneous class.  \\citet{PhillipsRelation} identified a relationship between the maximum B-band magnitude of an event and its rate of decline.  This ``brighter equals broader'' relationship has been extended to additional bands with templates from nearby events, allowing these events to be calibrated as an extension of the astronomical distance ladder (see \\citealt{jha2007} for a description).  This property, along with the brightness of SNeIa, which makes them visible over great distances, enables the use of SNeIa to probe the structure and expansion history of the universe, allowing studies of various cosmological models' parameters~\\citep{riessetal1998, perlmutter.aldering.ea:measurements, albetal2006, Kirshner09}, with recent work constraining cosmological parameters to within a few percent~\\citep{riessetal11, sullivanetal11}.  Recent observational studies of SNeIa have begun to correct for correlations of the brightness of a SNIa with properties of the host galaxy \\citep{ConleyEtAl11}.  Many SNIa observations are restricted to broadband photometry, so knowledge of host galaxy properties is correlated.  The inability to deconvolve these properties from each other is among the larger sources of uncertainty in cosmological constraints from SNeIa, so advancing the understanding of how brightness correlates with host galaxy properties may contribute significantly to reducing the uncertainties of cosmological parameters.  The brightness, and therefore ``broadness'', of a SNIa is determined principally by the amount of \\Ni{56} synthesized during the explosion. Observations report that SNeIa appear to have an intrinsic scatter of a few tenths of a magnitude after calibration, forcing a minimum uncertainty in any distances measured by using SNeIa as standardizable candles~\\citep{JacobyEtAl92, Kirshner09}.  An important goal of theoretical research into SNeIa, from the standpoint of cosmology, is to understand the sources of scatter and to identify potential systematic biases by studying the effects of various properties on the mechanism and nucleosynthetic yield of the SNIa.  The surrounding stellar population, the metallicity and mass of the progenitor, the thermodynamic state of the progenitor, the cooling and accretion history of the progenitor, and other parameters are known to affect the lightcurves of SNeIa; the role of these ``secondary'' parameters is the subject of considerable study~\\citep[e.g.,][]{Roepetal06_2, hoeetal2010}. Additionally, many of these effects may be interconnected in complex ways~\\citep{DomiHoefStra01, LesaffreEtAl06, townetal09}.  Observational campaigns are gathering information about SNeIa at an unprecedented rate.  \\citet{Scannapieco2005The-Type-Ia-Sup} and \\citet{MannucciEtAl06} showed that the delay time (elapsed time between star formation and the supernova event) data are best fit by a bimodal delay time distribution (DTD) with a prompt component that tracks less than 1 Gyr after star formation and a tardy component that occurs several Gyr later. \\citet{GallagherEtAl08} demonstrate a correlation between brighter SNeIa and shorter delay times, which they state is consistent with the bimodality described by~\\citeauthor{MannucciEtAl06}, but also with a continuous relation. \\citet{howelletal+09}, \\citet{NeillEtAl09} and \\citet{BrandtEtAl10} also find such a correlation between the delay time and brightness of a SNIa.  While the degeneracy of age and metallicity in observations could obscure these correlations, \\cite{howelletal+09} note that the scatter in brightness of this observed relation is unlikely to be explained by the effect of metallicity.  For this theoretical study, we adopt the model known as the single-degenerate paradigm.  This model assumes that a SNIa is the result of a thermonuclear disruption of a white dwarf (WD) in a mass-transferring binary system with either a main-sequence or red-giant companion star~\\citep[see] [and references therein]{branchetal1995, Fili97, hillebrandt.niemeyer:type, livio2000, roepke2006, LiEtAl2011, NugentEtAl2011, BloomEtAl2011}.  Recent observational evidence, however, suggests other progenitors such as the merging of two white dwarfs may explain many events~\\citep{scalzo:2010, Yuan:2010}.  In the single-degenerate scenario, the WD is formed when the primary star goes through a giant phase and expels a planetary nebula.  Once the primary becomes a WD, it is initially not in contact with the companion star, and it slowly cools as thermal energy is radiated away.  Once the companion star evolves and fills its Roche lobe, mass-transfer begins to carry low-mass elements from the envelope of the companion to the surface of the WD.  If the accretion rate exceeds $\\sim\\!\\!  10^{-7}$~\\Msol~yr$^{-1}$, the H-rich material can steadily burn~\\citep{NomotoEtAl07} and the WD gains mass, which heats and compresses the WD, driving up both the temperature and density in the core.  Once the temperature rises enough for carbon burning to begin, the core of the WD begins to convect; this is known as the ``simmering'' phase.  This simmering phase lasts on order of 10$^3$~yr, and ends when a flame is ignited, which occurs approximately when the eddy turnover time becomes shorter than the local nuclear runaway time.  Our initial models attempt to parameterize the WD at the end of the simmering phase, just at the beginning of the thermonuclear deflagration, which will in turn cause an explosion that will disrupt the entire WD in a SNIa.  The explosion mechanism we use (within the single-degenerate paradigm) is that of a deflagration to detonation transition (DDT). After ignition, the flame propagates as a subsonic deflagration for a while and then transitions to a supersonic detonation that rapidly consumes the star~\\citep{1986SvAL, woosley90, Khokhlov1991Delayed-detonat, hokowh95, HoefKhok96, KhokhlovEtAl97, NiemWoos97, hwt98, Niem99}. We describe the details of our implementation of this explosion mechanism below.  In the single-degenerate paradigm, a longer delay time can be explained by a longer elapsed time between the formation of the WD and the onset of accretion. During this period, the WD is in isolation and cools, hence the moniker the ``WD cooling time'' (\\tcool).  Following the cooling time is a period of accretion, during which the WD is compressed and heats, approaching the conditions for ignition of the thermonuclear runaway.  The decrease in temperature during the cooling time, which is determined by \\tcool, influences the density structure of the WD just prior to ignition, with a longer \\tcool\\ resulting in a higher central density when the core reaches the ignition temperature~\\citep{LesaffreEtAl06}.  Thus, a correlation between central density and the brightness of an event would suggest a correlation between delay time and the brightness of an event.  In this manuscript, we expand on our earlier investigation on the effect of \\tcool\\ on the brightness of the explosion.  In~\\citet{KruegerEtAl10} we reported that as the central density of the progenitor WD increases, the production of radioactive \\Ni{56} decreases due to increased neutronization rates, producing a dimmer event.  Using the results of \\citet{LesaffreEtAl06}, we related the WD central density to \\tcool\\ and were able to compare our results to the observations of \\citet{NeillEtAl09}.  Here we present additional details of our models; a statistical analysis of the results including the assessment of intrinsic scatter; the distribution of Fe-group elements within the remnant; and a potentially-observable effect to demonstrate the connection between age, progenitor central density, and brightness. We also revised our previously-reported trend in brightness with age to account for the main sequence evolution of the WD progenitor.  In \\secref{sec:method} we discuss the methodology of our suite of simulations, followed by details of the code we used to perform our simulations in \\secref{sec:code}.  We present the results of our simulations in \\secref{sec:results}, and discuss how these results compare with previous studies in \\secref{sec:discussion}.  \\secref{sec:conclusions} contains a brief summary and final conclusions. ",
        "conclusions": "\\label{sec:conclusions}  This paper builds on the results presented in \\citet{KruegerEtAl10}, giving more detail of the study and extending the analysis.  In that paper we showed that, in our 2-d simulations, a higher central density of the progenitor star does not impact the production of IGEs, but leads to greater neutronization, resulting in the production of less \\Ni{56}.  We also discussed the relation between the age of the progenitor and the central density (see, e.g., \\citealt{LesaffreEtAl06}), and the relation between the brightness and the mass of \\Ni{56} produced.  Thus the statement that a higher density leads to less \\Ni{56} is equivalent to the statement that an older progenitor will produce a dimmer SNIa.  In this work we expand on the discussion of \\citet{KruegerEtAl10} to give more detail of our models and to improve upon the age-brightness relation predicted by our simulations. In particular, we show that by adding a main-sequence lifetime to the cooling time our brightness-age relation is steeper, more closely matching the observed behavior of older SNeIa.  In comparing with other theoretical work, we see that the variation of \\Ni{56} mass with progenitor central density is not a settled question.  In this paper we further developed the idea that, due to the strong nonlinearities of the processes in SNeIa, a statistical study of an ensemble of SNIa simulations may be necessary to determine the true trends.  For our simulations, we find that 15 realizations (morphologies of the initial flame surface) are sufficient to characterize the mean trends from our models.  We find that the inner region (out to an enclosed mass of 0.8~--~1.0~$M_\\odot$) of the remnant is dominated by \\Ni{56}.  However, the stable (non-\\Ni{56}) IGEs tend to be in ``clumps'', instead of well-mixed throughout this region.  This may give rise to \\Ni{56} holes with little or no \\Ni{56}, depending on the line of sight through a SNIa remnant.  The outer region of the remnant will have more intermediate- and low-mass elements, as the burning becomes less efficient for lower densities.  As the central density increases, the mean \\Ni{56} mass fraction in the inner region drops (roughly 0.8 for $\\cdens = 1 \\times 10^9$~g~cm$^{-3}$ to roughly 0.6 for $\\cdens = 5 \\times 10^9$~g~cm$^{-3}$).  However, the extent of this region (in enclosed-mass space) does not significantly change.  Variations in the central density affect the sharpness of the edges of the stable IGE clumps: a higher central density leads to clumps of stable elements that are more sharply defined, as well as less mixing between the \\Ni{56} and the stable IGEs.  To better connect to observations, we discussed how to distinguish the relative ages of SNeIa with the same brightness (in other words, the relative initial central densities of SNeIa that produce the same mass of \\Ni{56}).  We found that, in our models, the best measure of the central density is the mass of stable IGEs, where higher central density progenitors produce more stable IGEs due to their greater rate of neutronization during the subsonic deflagration phase.  We found that a higher central density leads to a shorter deflagration phase. Since the rate of neutronization is significantly boosted, the total neutronization is greater at a higher central density despite there being less time to neutronize.  The time between the ignition of the first detonation and the cessation of burning is independent of central density.  As noted in \\citet{KruegerEtAl10} and described above, our choice for the DDT transition density led to an overproduction of \\Ni{56}.  In \\secref{sec:YieldFitting} we provided a recalibration of this overall brightness normalization to extrapolate our results to an expected average brightness.  Future work will be improved by a better choice of DDT density and we will report any quantitative changes to the trends reported here.  Due to the fundamentally 3-d nature of some of the phenomena in a SNIa (such as the turbulent velocities), we plan to extend this work by performing 3-d simulations.  Because 3-d simulations are much more computationally expensive than the corresponding 2-d simulations, a study with 3-d simulations will by necessity be constrained to a smaller number of simulations.  The choices will be motivated by this work and seek to span the parameter space explored here."
    },
    "1208/1208.6274_arXiv.txt": {
        "abstract": "The spin angular momentum $\\boldsymbol{S}$ of a supermassive black hole (\\sbh) precesses due to torques from orbiting stars, and the stellar orbits precess due to dragging of inertial frames by the spinning hole. We solve the coupled post-Newtonian equations describing the joint evolution of $\\boldsymbol{S}$ and the stellar angular momenta $\\boldsymbol{L}_j, j = 1\\ldots N$ in spherical, rotating nuclear star clusters. In the absence of gravitational interactions between the stars, two evolutionary modes are found: (1) nearly uniform precession of $\\boldsymbol{S}$ about the total angular momentum vector of the system; (2) damped precession, leading, in less than one precessional period, to alignment of $\\boldsymbol{S}$ with the angular momentum  of the rotating cluster. Beyond a certain distance from the \\sbh, the time scale for angular momentum changes due to gravitational encounters between the stars is shorter than spin-orbit precession times. We present a model, based on the Ornstein-Uhlenbeck equation, for the stochastic evolution of star clusters due to gravitational encounters and use it to evaluate the evolution of $\\boldsymbol{S}$ in nuclei where changes in the $\\boldsymbol{L}_j$ are due to frame dragging close to the \\sbh\\ and to  encounters farther out. Long-term evolution in this case is well described as uniform precession of the \\sbh\\ about the cluster's rotational axis, with an increasingly important stochastic contribution when \\sbh\\ masses are small. Spin precessional periods are predicted to be strongly dependent on nuclear properties, but typical values are $\\sim 10^7-10^8$ yr for low-mass \\sbhs\\ in dense nuclei, $\\sim 10^{8}-10^{10}$ yr for \\sbh\\ masses $\\sim 10^8\\msun$, and $\\sim 10^{10}-10^{11}$ yr for the most massive \\sbhs. We compare the evolution of \\sbh\\ spins in stellar nuclei to the case of torquing by an inclined, gaseous accretion disk. ",
        "introduction": "Introduction}  An accretion disk fed by gas whose angular momentum is misaligned with that of the central supermassive black hole  (\\sbh) will experience Lense-Thirring \\cite{LenseThirring1918} precession. Viscous torques near the \\sbh\\ align the gas with the \\sbh\\ equatorial plane \\cite{BardeenPetterson1975}; farther out, the gas remains inclined, producing a constant torque that causes the \\sbh\\ spin axis to precess. Such precession has been invoked as an explanation for changes in the direction of radio jets in active galaxies \\cite{BBR1980,Roos1988}. Continued accretion of gas from a misaligned plane will eventually reorient the \\sbh, although the time required for realignment is uncertain \\cite{NatarajanPringle1998}.  Accretion disks are believed to be associated with only a small fraction of \\sbhs. Here we consider the more generic, and perhaps simpler, case of a rotating \\sbh\\ embedded in a nuclear cluster of stars or stellar remnants. If  the  cluster has a net  angular momentum that is misaligned with the \\sbh\\ spin, a mutual torque will be exerted between stars and \\sbh, even if the {\\it spatial} distribution of the stars is precisely spherical.  In the simplest such model, the stars move independently of each other. % Differential precession (``phase mixing'') will nevertheless cause stellar orbits near the \\sbh\\ to distribute their angular momentum vectors $\\boldsymbol{L}_j$ uniformly about the spin $\\boldsymbol{S}$,  decreasing the torque that they exert on the hole. The angular momentum associated with stars farther out can remain misaligned, leading to a forced precession of the \\sbh, similar to what occurs in the case of misaligned accretion disks.  By solving the coupled post-Newtonian equations describing a spinning \\sbh\\ and a rotating cluster, we verify that such an outcome is possible, at least starting  from certain initial conditions. However we find a second evolutionary mode as well, in which differential precession causes the inner system to reach alignment with the total (spin plus orbital) angular momentum, resulting in a steady state with no subsequent precession of the hole.  Stars also interact with each other gravitationally; these encounters lead to changes in stellar  angular momenta, on time scales that can be short compared with Lense-Thirring  times. Unlike changes due to frame-dragging, evolution of the $\\boldsymbol{L}_j$ due to encounters is essentially random. There is a region near the \\sbh, the ``sphere of rotational influence,'' in which encounter times are long compared with frame-dragging times. Within this region, stellar orbits precess uniformly, while outside of it, changes in the $\\boldsymbol{L}_j$ are due primarly to encounters and are random. The size of this sphere varies from $\\sim 10^{-3}$ pc in the nuclei of galaxies like the Milky Way, to $\\sim 10^1$ pc in nuclei containing the most massive \\sbhs. We develop a stochastic model for the evolution of $\\boldsymbol{S}$ that includes the effects of encounters on the  $\\boldsymbol{L}_j$. In this model, net alignment of the stellar angular momena with the \\sbh\\ spin is less efficient, and the \\sbh\\ typically continues to precess about the mean $\\boldsymbol{L}$ of the stellar cluster, although its instantaneous precession rate can vary stochastically due to the stochastically changing $\\boldsymbol{L}_j$.  Evolution of \\sbh\\ spins to due torquing from stars has many parallels with evolution due to torquing from an accretion disk, surprisingly so given that one process is energy-conserving and the other is dissipative. We compare and contrast the two sorts of evolution in the ``Discussion'' section, where we also summarize observational and theoretical evidence for nuclear rotation, and discuss the implications of our results for the experimental determination of black hole spins.  Throughout this paper we ignore the contribution of stellar captures to the evolution of $\\boldsymbol{S}$. ",
        "conclusions": "\\label{Section:Conclusions}  \\noindent 1. In a galactic nucleus containing a spinning supermassive black hole (\\sbh), frame dragging results in mutual torques between the stellar orbits and the \\sbh. The result is precession of both the \\sbh\\ spin, $\\boldsymbol{S}$, and the angular momentum vectors, $\\boldsymbol{L}_j$, of the individual stellar orbits, with $\\boldsymbol{S}+\\boldsymbol{L}_\\mathrm{tot} = \\boldsymbol{S} + \\sum_j\\boldsymbol{L}_j$ conserved. For stars at a single distance from the \\sbh, the controlling parameter is the ratio between $S$ and $L_\\mathrm{tot}$. If $S\\gg L_\\mathrm{tot}$, stellar orbits precess about the nearly fixed $\\boldsymbol{S}$ with the Lense-Thirring period; while if $L_\\mathrm{tot}\\gg S$, $\\boldsymbol{S}$  precesses about the nearly fixed $\\boldsymbol{L}_\\mathrm{tot}$ with a period that is shorter by a factor $S/L_\\mathrm{tot}$. The inner parsec of the Milky Way is known to contain stellar subsystems having $L_\\mathrm{tot}\\approx S$.  \\smallskip  \\noindent 2. Ignoring interactions between the stars, solutions of the coupled equations for $\\dot{\\boldsymbol S}$ and $\\dot{\\boldsymbol L}_{j=1,\\ldots,N}$ in spherical nuclei reveal two evolutionary modes in the case that $L_\\mathrm{tot}>S$: continued precession of $\\boldsymbol{S}$ about $\\boldsymbol{L}_\\mathrm{tot}$; or damped precession, in which $\\boldsymbol{S}$ and $\\boldsymbol{L}_\\mathrm{tot}$ come into nearly complete alignment after one precessional period of the \\sbh. Even in the first  mode, differential precession of orbits near the \\sbh\\ causes their net angular momentum to align with $\\boldsymbol{S}$, reducing the torque that they exert on the \\sbh. Subsequent precession of the \\sbh\\ is  driven by torques from stars at $r\\gap r_\\mathrm{L}$, where $r_\\mathrm{L}$ is the radius enclosing a net angular momentum equal to $S$.  \\smallskip  \\noindent 3. Newtonian interactions between stars can change their orbital angular momenta in a time shorter than Lense-Thirring precessional times. We define the ``radius of rotational influence,'' $a_\\mathrm{K}$, around a Kerr \\sbh\\ as the radius inside of which torques due to  frame dragging act more quickly than torques from the other stars. Typical values for this radius are $\\sim 10^{-3}$ parsecs in dense nuclei like that of the Milky Way, increasing to $\\sim 10^0-10^1$ parsecs in nuclei containing the most massive \\sbhs. The angular momentum associated with stars in this ``collisionless'' region near the \\sbh\\ is likely to be much smaller than $S$ in nuclei of the smallest galaxies but may be comparable to $S$ in massive galaxies.  \\smallskip \\noindent 4. Interaction between stars at $r>a_\\mathrm{K}$ leaves the total angular momentum of these stars unchanged, but results in random fluctuations of the individual $\\boldsymbol{L}_j$ and hence in the torque which they exert on the \\sbh. We develop a stochastic model, based on the Ornstein-Uhlenbeck equation, for the torque exerted by these stars and verify it by comparison with high-accuracy $N$-body simulations. We argue that $d\\boldsymbol{S}/dt$ can be approximated as the sum of two terms: deterministic torques exerted by stars inside $a_\\mathrm{K}$, whose angular momenta evolve solely in response to frame-dragging; and a stochastically-fluctuating torque due to stars outside $a_\\mathrm{K}$.  \\smallskip \\noindent 5. Examples of stochastic evolution of $\\boldsymbol{S}$ are presented for various nuclear models. Typical evolution consists of sustained precession, with periods that are highly dependent on nuclear parameters, but which are expected to increase with increasing $\\mh$: likely periods are $\\sim 10^7-10^8$ yr for low-mass \\sbhs\\ in dense nuclei, $\\sim 10^{8}-10^{10}$ yr for \\sbh\\ with masses $\\sim 10^8\\msun$, and $\\sim 10^{10}-10^{11}$ yr for the most massive \\sbhs."
    },
    "1208/1208.4359.txt": {
        "abstract": "Many hot subdwarf stars show composite spectral energy distributions indicative of cool main sequence companions. Binary population synthesis (BPS) models demonstrate such systems can be formed via Roche lobe overflow or common envelope evolution but disagree on whether the resulting orbital periods will be long (years) or short (days). Few studies have been carried out to assess the orbital parameters of these spectroscopic composite binaries; current observations suggest the periods are long. To help address this problem, we selected fifteen moderately-bright (V$\\sim$13) hot subdwarfs with F--K dwarf companions and monitored their radial velocities (RVs) from January 2005 to July 2008 using the bench--mounted Medium Resolution Spectrograph on the Hobby--Eberly Telescope (HET). Here we describe the details of our observing, reduction, and analysis techniques and present preliminary results for all targets.  By combining the HET data with recent observations from the Mercator telescope, we are able to calculate precise orbital solutions for three systems using more than 6 years of observations.  We also present an up--to--date period histogram for all known hot subdwarf binaries, which suggests those with F--K main sequence companions tend to have orbital periods on the order of several years.  Such long periods challenge the predictions of conventional BPS models, although a larger sample is needed for a thorough assessment of the models' predictive success. Lastly, one of our targets has an eccentric orbit, implying some composite--spectrum systems might have formerly been hierarchical triple systems, in which the inner binary merged to create the hot subdwarf. ",
        "introduction": "\\label{sec:intro}  Binarity plays an essential role in the story of the hot subdwarf B (sdB) stars, one of the most enigmatic stages of stellar evolution.  The sdBs dominate surveys of faint blue objects and are found in almost all Galactic stellar populations. They are the field counterparts of the Extreme Horizontal Branch stars in globular clusters, and their location in the H--R diagram corresponds to stellar models with He--burning cores and extremely thin H envelopes \\citep{heb86}. Presumably, the sdB progenitors were stripped of almost all their surface hydrogen while on or near the red giant branch (RGB), igniting helium burning in their cores at nearly the same time. The question is how this can happen, since most Population I stars do not lose their entire envelopes, and remain in the red clump, not far from the RGB, during core He burning.  \\begin{table*} \\centering \\caption{RV Monitoring Targets} \\scriptsize \\begin{center} \\leavevmode \\begin{tabular}{lllllll} \\hline \\hline Target & Sp. Type$^a$ &RA & Dec & V & Alternate &Comments\\\\ &                  & [J2000] & [J2000] & & Name & \\\\ \\hline PG 0039+049 &  sdB+G2V \t&  00:42:06.1\t&+05:09:24 \t&   12.9\t&  PB 6107 & \\\\%F9-G4 PG 0110+262 & sdB+G0V   \t\t& 01:13:14.9   \t&+26:27:31  \t&   12.9   \t&    & \\\\ %not best, F8-G2 PB 8783 & sdOV$^{b}$+F4V  \t\t& 01:23:43.2  \t&  -05:05:45   \t&    12.3    & EO Cet &  pulsating hot subdwarf$^5$ \\\\  %F2-6 PHL 1079 & sdB+G7V\t\t& 01:38:27.1  \t& +03:39:39   \t& 13.4  \t&   &          \\\\ %good, G5-9 PG 0232+095 & sdB+G1V \t\t& 02:35:12.0   \t&+09:45:38\t&  12.5    \t&      &\\\\ %nobest chi-sqr F8-G4 PG 1040+234 & sdB+G1V \t&10:43:39.3  \t& +23:09:07   \t& 13.4   \t& TON 1281& resolved$^{1}$    \\\\  %not best G PG 1104+243 &  sdOB+G2V \t& 11:07:26.3   \t&+24:03:12   \t& 11.3  \t&  &  MS orbit published$^4$\\\\   %F9-G3 PG 1206+165 & sdB+G9V$^c$\t& 12:09:16.7   \t&+16:11:56   \t& 13.8    \t&  PB 3854 & \\\\  %hardly any companions lines! PG 1253+284 &sdB+F8V\t& 12:56:04.9   \t&+28:07:20 \t& 12.7   \t& TON 139    &sdB orbit published$^{3}$; resolved$^{1}$; presumed triple system$^{1,3}$ \\\\ %F5-G2 PG 1317+123 &  sdO+G1V\t& 13:19:53.6   \t&+12:03:59   \t& 11.3     \t & Feige 80  & MS orbit published$^4$  \\\\ %F8-G3 PG 1338+611 &  sdB+G4V \t&13:40:14.7   \t&+60:52:48      \t&11.4\t&   Feige 87   & MS orbit published$^4$    \\\\%G3-7 PG 1449+653 & sdB+G5  \t&14:50:36.1   \t&+65:05:53   \t&13.6\t&   &    \\\\  % messy:early G PG 1629+081 &  sdOB+K5V$^c$ \t& 16:32:01.4   \t& +07:59:40   \t& 12.8\t&   & resolved$^{2}$  \\\\   %no lines from cool companion, which wasn't on fiber PG 1701+359 & sdB+G8\t& 17:03:21.6 \t&  +35:48:49   \t& 13.2\t&   &    \\\\  %G6-K0 PG 1718+519 & sdB+G4V\t&17:19:45.5  \t& +51:52:10 \t& 13.7\t&     &  resolved$^{1}$ \\\\   %G2-6 \\hline \\multicolumn{7}{l}{$^{a}$classifications from this work (see \\S \\ref{subsec:classifications}) unless otherwise noted}\\\\ \\multicolumn{7}{l}{$^{b}$sdO status of the hot subdwarf recently confirmed by \\citet{ost12b}}\\\\ \\multicolumn{7}{l}{$^{c}$companion classification from \\cite{staphd}}\\\\ \\multicolumn{7}{l}{\\textbf{References}: $^{1}$\\citet{heb02}; $^{2}$\\citet{ost05}; $^{3}$\\citet{cop11}; $^4$\\citet{ost12}; $^5$\\citet{koe97} }\\\\ \\end{tabular} \\end{center} \\label{tab:targets} \\end{table*}  Stochastic mass loss from single stars \\citep{dcr96} or other single--star scenarios to explain this are \\textit{ad hoc}, so most recent effort has concentrated on mass loss resulting from interactions in a close binary star system, following the early exploration by \\citet{men76}. \\citet{han02,han03}, for example, described five channels of binary star evolution that can lead to the formation of hot subdwarfs: the `first' and `second' stable Roche--lobe overflow (RLOF) scenarios, the `first' and `second' common envelope (CE) channels, and the merger of two He-core white dwarfs (WDs).  Their CE scenarios produce close binaries with orbital periods from hours to days; the sdB companions are generally late-type (G/K/M) main sequence stars or white dwarfs.  In contrast, their stable RLOF channels lead to systems with early--type (B/A/F) main sequence or late--type giant companions and long orbital periods ($P \\sim$\\, few $\\times\\ 100$ d).  An alternative scenario employing the $\\gamma$--formalism has been put forth by \\citet{nel00,nel01a} and \\citet{nel10}, whose CE channels can produce systems with main sequence companions and periods on the order of years.  To date, orbital parameters have been measured for approximately 150 sdB binary systems; the overwhelming majority of these are close binaries with $P < 15$ d (see Table A.1 of \\citealt{muchfuss} for a recent summary).  The companions are mostly WDs and M dwarfs, essentially invisible in the glare of the sdB.  With a suitable choice of tuning parameters, theoretical models can reproduce the observed distribution of sdB masses and orbital periods in these systems reasonably well \\citep{max01,cop11}.  The recent binary population synthesis (BPS) study by \\citet{cla12}, however, shows that numerous parameter sets can reproduce this observed subpopulation; thus, \\textit{additional constraints are needed to better tune BPS codes}.  Very little attention has been paid to the sdB+F/G/K binaries, in spite of the nearly equal flux contributions from the two components in such systems. The relative dearth of published orbital parameters implies their periods are quite long.  More than a decade ago, \\citet{gre01} hinted that sdB binaries with composite spectra tend to have longer orbital periods.  They reported an average around 3-4 years, although no specific systems were named. Since that brief discussion, only two other studies have claimed the detection of orbital periods in composite--spectrum binaries, both of which were published very recently. \\citet{ost12} (hereafter \\ostrv) discuss a survey conducted with the High Efficiency and Resolution Mercator Echelle Spectrograph (HERMES, see \\citealt{ras11}) mounted at the Mercator Telescope on La Palma, Spain and present preliminary results showing orbital periods longer than $\\sim$ 500 d for at least eight binaries. In addition to that work, \\citet{dec11} report a period of 760 d for PG 1018$-$047.  Orbits with such long periods are sufficiently large to have once accommodated an inner binary, so it is possible that these systems were formerly hierarchical triple systems, in which case the subdwarf could have formed from the merger of the inner binary.  \\citet{cla11}, for example, recently proposed a scenario in which the merger of a He--core WD and a low--mass main sequence star could eventually form an sdB.  If such a binary were in orbit with a dwarf companion, the merger would leave behind an sdB and a main sequence star that had no part in the formation of the sdB (aside from potentially advancing the merger process via the Kozai mechanism).  Unlike stable RLOF-- and CE--produced systems, which should have nearly circular orbits, no limitations exist on the eccentricities of these binaries other than a requirement that the periastron separation not be too small.  Thus, it is imperative to constrain the orbital geometry in addition to the period and radial velocity (RV) amplitudes.  %It is important to characterize all subpopulations of sdB binaries since they %help constrain the parameterizations in BPS codes, which are used throughout %astrophysics to model a wide variety of exotic phenomena, including LMXBs, %black hole binaries, SN Ia, and CV\u00eas, in addition to sdB stars. Since hot %subdwarfs %are believed to be the main source of the \u0096UV-upturn\u0094 in old stellar %populations %(see e.g., O\u00eaConnell 1999; Brown et al. 1997, 2000; Han, Podsiadlowski \\& %Lynas-Gray 2007), %improvements in BPS modeling help us better understand the source of UV light %from stellar populations, which becomes more important as the rest-frame UV %is probed at higher redshifts.  Here we describe a survey carried out with the Hobby--Eberly Telescope (HET) over $\\approx 3.5$ years to measure the line--of--sight accelerations of a sample of sdBs with composite spectra.  We discuss our target selection in \\S \\ref{sec:targets} and observing strategies in \\S \\ref{sec:obs}.  The details of our analysis techniques for determining velocities and fitting the RV curves are given in \\S \\ref{sec:analysis}.  Preliminary results and the current status of our program are presented in \\S \\ref{sec:results}.  After combining the HET data with recent measurements from \\ostrv , we report precise orbital solutions for three of the targets and discuss them in greater detail.  Our initial results confirm early suggestions that the orbital periods of many sdB+F/G/K systems are long, on the order of several years.  We present in \\S \\ref{sec:histogram} an updated orbital period histogram for all measured sdB binaries, which shows a possible gap in the period distribution centered near $P=70$ days, and we briefly discuss the puzzle that it presents and possible resolutions.  Finally, we summarize our results in \\S \\ref{sec:conclusions}.  %[*For thought: do we detect any %unresolved but non-accelerated systems??? if FGK unrelated to sdB %formation, shouldn't the mean period be ~280 years... Duquennoy \\& Mayor %'91; 50 AU at ~kpc is likely unresolved... Unless Kozai mechanism is %essential...]  %%%%%%%%%%%%%% [---  TARGET SELECTION--] %%%%%%%%%%%%%%%% %\\vspace{5mm} ",
        "conclusions": "\\label{sec:conclusions} We have described an ongoing program to determine the orbital parameters of hot subdwarfs in binaries with F--K type main sequence companions. Preliminary results show most of our targets have orbital periods on the order of years (rather than days), in agreement with the recent study by \\ostrv\\ and the findings of \\citet{gre01} from more than a decade ago. Precise orbital solutions cannot be provided for most of our systems at this time, owing to insufficient phase coverage.  Nonetheless, the preliminary findings imply that the assumptions inherent in some subdwarf formation scenarios should be revised, which probably requires both re--tuning the parameters used in BPS models and accounting for at least some of the long--period systems as former hierarchical triples.  % The following text used to be at end of first paragraph, after \"hierarchical triples\": % %It is \\textit{imperative} that observers continue probing %these longer periods to further constrain viable binary evolutionary pathways. % We are currently conducting a follow--up program %with HET/HRS to determine the orbital periods, velocities, and %eccentricities of the unsolved targets in our survey.  For the three systems also observed by \\ostrv, we calculated precise orbital parameters from nearly seven years of RV measurements.  Mass ratios were derived from the orbital RVs and are consistent with our spectral classifications for the cool companion and canonical--mass hot subdwarfs for two of the three systems.  In the case of the third, PG 1104+243, the derived subdwarf mass exceeds the range classically expected for sdBs that evolved through the stable RLOF channel, implying the subdwarf either formed from a WD merger (resulting in a larger--than--canonical mass), or the companion has a mass different from what we infer from its spectral type.  A long--term goal for observers is to determine the joint distribution of orbital periods and eccentricities in long--period sdB binary systems. At this time, we cannot rule out eccentric orbit solutions for PG 1104+243 or PG 1317+132.  Determining the eccentricity of PG 1317+123 will prove especially difficult due to the presence of rapid RV variations, which might originate from star spots on the cool companion.  We find a statistically significant eccentricity of $e=0.15 \\pm 0.02$ for PG 1338+061 and reject the circular--orbit hypothesis in this case.  The long--period binary PG 1018$-$047 also shows evidence in favor of a non--circular orbit \\citep{dec11}, although this result needs confirmation. Such departures from $e=0$, taken at face value, challenge the prevailing notion that strong binary interactions, involving the filling of the sdB progenitor's Roche lobe, are always involved in the formation of hot subdwarf stars. If each of these binaries was (or is) a hierarchical triple system, then the cool companion that we observe probably did not participate directly in the formation of the hot subdwarf, in which case its orbit would not have been circularized. The hot subdwarf in this picture could have formed from the merger of two He--core WDs \\citep{han02,han03} or a He--core WD and a MS star \\citep{cla11}. Measurements of the rotational velocities of sdBs in long--period systems might shed light on their formation by identifying the type of merger. On the other hand, if the system has always been a binary, a strong, tidally--enhanced wind \\citep{tou88} could have assisted the sdB progenitor in ejecting its envelope on the RGB before it expanded enough to fill its Roche lobe, thereby preserving some level of eccentricity in the binary.  Even in the case the Roche lobe was filled previously, non--steady mass loss, the action of a distant third companion, or interactions with a circumbinary disc could have potentially added eccentricity to the orbit, although these scenarios seem less likely.  In light of the updated period histogram for sdB binaries (Fig. \\ref{fig:period_histogram}), we find it \\textit{imperative} that observers continue to study both short-- and long--period systems with all possible types of companions, to arrive at a clear understanding of the formation channels for hot subdwarf stars.  Thus it is important to give attention to issues of completeness and bias in our knowledge of binary orbits (or otherwise) for sdB stars.  An important first step is to establish a more complete catalogue or finding list, as for example in the effort described by \\citet{gir12}, who search for photometrically composite objects and use decomposition of the spectral energy distribution to identify binaries containing hot subdwarfs.  We note that one of the selection criteria of \\citet{gir12} is that the binary is unresolved; this by itself does not indicate whether the binary is close enough to be tidally interacting, so follow--up studies are needed. Eventually the goal is to derive full orbital elements for an unbiased sample of any interacting binaries that are found in these surveys.  Meanwhile, work continues to characterize the orbits of sdB binaries that are already known.  We are currently conducting a follow--up program with HET/HRS to determine the orbital periods, velocities, and eccentricities of the unsolved targets in our survey."
    },
    "1208/1208.4101_arXiv.txt": {
        "abstract": "We investigate the capability of pulsar timing arrays (PTAs) as a probe of primordial black holes (PBHs), which might constitute the Galactic dark matter. A PBH passing nearby the Earth or a pulsar gives an impulse acceleration and induces residuals on otherwise orderly pulsar timing data. We show that  the timing residuals induced at pulsars are optimal for searching heavier PBHs than those at the Earth, and the two probes are highly complemental. Future facilities like SKA could detect PBHs with masses around $\\sim 10^{22\\mbox{-}28} \\mathrm{g}$ even if only a small fraction ($\\lesssim 1\\% $) of the Galactic dark matter consists of these PBHs. ",
        "introduction": "A significant fraction of matter in our Galaxy is considered to be occupied by dark matter, but  its nature is poorly understood at present \\citep{2005PhR...405..279B}. Primordial black holes (PBHs)  are an interesting astrophysical candidate of dark matter, and various observational constraints  have been posed on their allowed mass range \\citep{2010PhRvD..81j4019C,2010RAA....10..495K}. However, the current constraints  in the mass range $10^{20}{\\rm g}<M_\\mathrm{PBH}<10^{27} {\\rm g}$  remain relatively  weak.   To examine PBHs in this range, \\cite{2007ApJ...659L..33S} proposed an observational technique to directly probe gravitational interaction between the solar system and a nearby PBH. Around their close approach,  the Earth receives an impulse of acceleration whose profile depends on the mass, distance, and velocity of the PBH. Given the local mass density of dark matter, the expected magnitude of the acceleration is very small, but high-precision  measurements of  pulsar timing  could allow us to detect the weak impulse signal (see also \\cite{2004PhRvD..70f3512S,2009PhRvL.102p1101S,Griest:2011}).   Roughly speaking, in the present context,  the pulsar timing analysis can be essentially regarded as measurement of the arrival times of radio pulses that were emitted by a pulsar and received on the Earth.  The local accelerations of both of the pulsar and the Earth are encoded in the modulation of the time-of-arrival (TOA) data, as a simple linear combination of two separate terms for the two masses.  We call them by the pulsar term and the Earth term respectively.    When observing multiple pulsars for  PBH search, the Earth terms are commonly excited by the  acceleration of the Earth, and have coherent structure among TOA data of different pulsars. Therefore, we can statistically amplify the weak impulse signature on the Earth by using a pulsar timing array (PTA) and  effectively reducing the timing noises.  The underlying statistical approach here is similar to that for detecting gravitational waves (GWs) whose effect on the TOA data can be also expressed with  two terms induced  separately  at the pulsar and the Earth \\citep{1978SvA....22...36S,1979ApJ...234.1100D,1983ApJ...265L..39H}. For the  GW signals, we call them by the pulsar GW term and the Earth GW term, in order to distinguish them from the two acceleration terms.  The traditional target of PTAs is the Earth GW terms.    Recently, it has been actively discussed that we  might utilize the pulsar GW terms for analyzing GW sources such as  merging super-massive black hole binaries \\citep{2004ApJ...606..799J,2010arXiv1008.1782C,2011MNRAS.414.3251L,Ellis:2012}. To this end, we need sufficiently stable millisecond pulsars (MSPs), and such preferred systems might be discovered   with future observational facilities.  With the pulsar GW terms and known distances to the individual pulsars, our information content on GWs at the nano-Hertz band can be greatly increased, compared with that extracted from  the Earth GW terms alone.   Meanwhile,  if the timing noises of individual pulsars are small, we can also probe PBHs around the pulsars through the pulsar (acceleration) terms, in addition to PBHs close to the solar system. Then, the pulsar terms can largely widen our survey volume, and our sensitivity for the direct PBH search would become better than the previous estimation only with the Earth terms \\citep{2007ApJ...659L..33S}. In  this paper, we study this issue with special attention to the differences between the roles of two terms for the PBH search with PTAs, namely, the independent and numerous pulsar terms and the coherent Earth terms. We find that these terms have both advantages and disadvantages for PBH search, and work in a complementary way.      % ",
        "conclusions": "Our order-of-magnitude estimation has shown that we might probe PBHs or constrain their abundance with future PTAs. Here let us discuss  issues related to  the PBH search.  \\subsection{Data analysis and competing noises} So far we have simply evaluated the signal-to-noise ratios of fly-by events, assuming  white noise spectra for timing noises of MSPs.  But  we should pay much attention to whether PBH signals can be clearly identified in their  TOA data.  The actual timing analysis typically proceeds as follows (see {\\it e.g.} \\cite{2006MNRAS.369..655H, 2006MNRAS.372.1549E} for the detail). First, the measured TOAs are  converted to the pulse emission times in the presumed reference frame of each MSP. In this step, various effects are estimated, such as the propagation delays, the ephemerides, and the coordinate transformation from the Earth to the pulsar. The derived time of emission is fitted by a timing model including a pulse frequency and its time derivative. The residual is divided into a white noise and un-modeled systematic noises which could include the signal we are seeking. A key in our approach to enlarge the detectable range of PBH mass is to consult both the coherent Earth terms and the individual pulsar terms. These two are assumed to be separated.  Here, we should be aware of the risk to mistake PBH signals for other effects while processing the TOA data, which could also lead to misestimation of the parameters of each pulsar. To prevent this, more detailed modeling of the PBH signal is required.   %    With a PTA, we have a large number of independent pulsar terms. The impulse accelerations by fly-by PBHs appear in the pulsar terms as sparse and isolated signals with finite durations. Therefore, a long-term observational campaign  would be quite helpful to select sufficiently stable MSPs and examine  the characteristic time profiles of the PBH signals.   As mentioned above,  we have assumed that the ordinary timing noise in each TOA is white, $\\sigma_\\mathrm{WN} \\propto T^{0}$ ($T$: the observational time span). As for non-recycled pulsars, there have been confirmed secular red noises, $\\sigma_\\mathrm{RN} \\propto T^{2 \\pm 0.2}$ \\citep{2010ApJ...725.1607S}. Although most of the observed MSPs only have upper limits on the amplitude of the red noise, it might be identified by future PTAs to be obstacle in the PBH search.  \\subsection{Competing signals} Once a higher sensitivity as we have considered is realized, other possible signals also have to be opened up for discussion. They can be effective noises for PBH search and vice versa.  The most promising target of future PTAs is the GW background from merging supermassive black holes. The typical magnitude of the estimated dimensionless strain is  $h \\sim 10^{-15.5} \\times (f/0.1\\mathrm{yr}^{-1})^{-2/3}$ \\citep{2003ApJ...583..616J, 2003ApJ...590..691W}. This corresponds to the amplitude of the timing residual of $s_{f,\\mathrm{GW}} \\sim 4.1 \\mathrm{ns} \\times (f/0.1\\mathrm{yr}^{-1})^{-5/3}$ (for the translation from $h$ to $s_{f,\\mathrm{GW}}$, see {\\it e.g.} \\cite{2009MNRAS.394.1945H}), which can be comparable to that induced by PBHs (see Eq.(\\ref{timing_residual})). Given the effective sensitivities of the pulsar and the Earth terms, the latter would be more vulnerable to the existence of the GW background.  But, using the angular patterns of the Earth terms,  we can, in principle, separate the two competing signals, as follows \\citep{2007ApJ...659L..33S}. % A PBH signature in the Earth terms  will have a dipole pattern ($l=1$) whose direction is determined by the acceleration vector. On the other hand, the GW background will have multiple modes starting from the quadrupole ($l=2$).  One may suspect that floating planets \\citep{Sumi:2011} could also become an effective noise because of the similar masses to PBHs we are interested in. There are expected to be a few times more floating planets than stars. The anticipated encounter rate with a pulsar is $\\sim 10^{-(7\\mbox{-}8)} \\mathrm{yr}^{-1} \\times (b/740\\mathrm{AU})^2 $, which is much smaller than PBH constituting a dominant fraction of dark matter (see Eq.(\\ref{event_rate})). Thus, we conclude that floating planets would not be problematic for our approach.  \\subsection{Possible refinements} Finally let us discuss possible improvements for our estimation of the detectable region of the PBH parameters.  In Fig.2 and Fig.3, we have fixed $V = 350 \\mathrm{km/s}$, $|\\cos \\theta | = 0.58$, and $\\rho_\\mathrm{DM} = 0.011 M_\\mathrm{\\odot} \\mathrm{pc}^{-3}$. These values correspond to the observed rms velocity for halo dark matter relative to the solar system, the mean projection factor for isotropic scatterings, and the local dark matter density near the solar system, respectively. These treatments should be regarded as a zeroth order approximation, and  we can refine the present analysis in a  more realistic manner.   For example, while the motions of the MSPs would be typically  close to the Galactic rotation, those of PBHs would be more isotropic. Then  the mean values for the projection factor $\\cos \\theta$ and the relative velocity $V$ would depend on the Galactic position of each target MSP. In  addition, the densities of the PBHs and the MSPs would also depend on the Galactic position.  These would affect the estimation of the  event rate and the strategy for data analysis.   %   Future facilities like SKA could determined the distance of MSPs up to $13 \\mathrm{kpc}$ with $\\lesssim 20 \\%$ accuracy \\citep{Smits:2011}. Multiple detections using the pulsar terms could potentially give us more detailed informations of the PBHs, such as their density and velocity distributions in the whole Galaxy, which can not be reached only using the Earth terms."
    },
    "1208/1208.4337_arXiv.txt": {
        "abstract": "Results of a statistical analysis of solar granulation are presented. A data set of 36 images of a quiet Sun area on the solar disk center was used. The data were obtained with the 1.6~m clear aperture New Solar Telescope (NST) at Big Bear Solar Observatory (BBSO) and with a broad-band filter centered at the TiO (705.7~nm) spectral line. The very high spatial resolution of the data (diffraction limit of 77~km and pixel scale of 0.$''$0375) augmented by the very high image contrast (15.5$\\pm$0.6\\%) allowed us to detect for the first time a distinct subpopulation of mini-granular structures. These structures are dominant on spatial scales below 600~km. Their size is distributed as a power law with an index of -1.8 (which is close to the Kolmogorov's -5/3 law) and no predominant scale. The regular granules display a Gaussian (normal) size distribution with a mean diameter of 1050 km. Mini-granular structures contribute significantly to the total granular area. They are predominantly confined to the wide dark lanes between regular granules and often form chains and clusters, but different from magnetic bright points. A multi-fractality test reveals that the structures smaller than 600~km represent a multi-fractal, whereas on larger scales the granulation pattern shows no multi-fractality and can be considered as a Gaussian random field. The origin, properties and role of the newly discovered population of mini-granular structures in the solar magneto-convection are yet to be explored. ",
        "introduction": "Solar granulation is visible in broadband filter images as a cellular pattern of bright features separated by dark lanes and is regarded to be a manifestation of convection in the outermost layers of the solar convective zone (e.g., Nordlund et al. 2009). Physical properties of the granulation are relevant to energy transfer from the sub-photospheric convective layers into the photosphere. This causes perpetual research interest in solar granulation,  which elevates any time when a newer, higher resolution solar instrument comes on line (e.g., Roudier \\& Muller 1986; Schrijver et al. 1997; S{\\'a}nchez Cuberes et al, 2000; Danilovic et al. 2008; Yu et al. 2011).  Roudier and Muller (1986) analyzed 0$''$.25 resolution white-light photographs of the quiet Sun granulation obtained with the Pic-du-Midi 50-cm refractor. They found a gradual decrease in granular size distribution function that implies absence of any dominant spatial scale. At the same time, these authors found that granules of about 1000 km in size contribute most to the total granular area. They referred to this 1000 km scale as a ``dominant scale'' and associated it with the characteristic thickness of the top layer of the convective zone.  Schrijver et al. (1997) utilized quiet Sun granulation images from the Royal Swedish Observatory at La Palma with an effective spatial resolution of 0$''$.4. They also reported that the granular cell size distribution is compatible with that reported by Roudier and Muller.  A key property of images of granulation is the root-mean-square of the intensity fluctuations, $\\Delta I_{rms}$ (Roudier \\& Muller, 1986; S{\\'a}nchez Cuberes et al, 2000; Danilovic et al. 2008; Yu et al. 2011): \\begin{equation} \\Delta I_{rms}=(\\Sigma(I-I_0)^2/NI_0^2)^{1/2}, \\label{rms} \\end{equation} where $I_0$ is the mean intensity of the image and $N$ is a number of data points, and the sum is taken over all image data points. Recently this parameter was referred to as the {\\it granulation contrast} (see, e.g., Danilovich et al. 2008). (Note that Roudier and Muller (1986) reserved the term ``contrast'' for a different parameter.) In this paper, we will refer to  $\\Delta I_{rms}$ calculated from Eq. (\\ref{rms}) as the granulation contrast.  Danilovic et al. (2008) compared the granulation contrast derived from MHD-simulations to Hinode/SP data (continuum images at 630~nm, Tsuneta et al. 2008). They reported a granulation contrast from the simulation data about 14-15\\%, which was found consistent with the observed granulation contrast of 7\\% , after appropriate degradation of the simulated data. According to S{\\'a}nchez Cuberes et al. (2000), raw ground-based and balloon observations do not produce a granulation contrast, $\\Delta I_{rms}$, higher than 9.4\\%, which was obtained by Rodriguez Hidalgo et al (1992) with the 0.5~m Swedish Vacuum Solar Tower (La Palma) at a wavelength of 468.6 nm.    In this study, we present the results of analysis of the granulation contrast,  size distributions, and multi-fractal properties of solar granulation using quiet-Sun data near the disk center obtained with the 1.6 m New Solar Telescope (NST, Goode et al. 2010a,b) at the Big Bear Solar Observatory (BBSO). ",
        "conclusions": "Solar granulation data measured with the NST's broad-band TiO imager display a very high granulation contrast (r.m.s. brightness fluctuation, in the terminology of Roudier \\& Muller, 1986) of 15.5$\\pm$0.6\\%, which agrees very well with the the 14-16\\% contrast for the TiO spectral line inferred from the numerical simulations using the STOPRO radiative transfer code (Kitiashvili et al. 2012). The NST contrast also comparable with the 14-15\\% contrast inferred from 3D MHD model data for the Hinode 630.25~nm spectral line (Danilovic et al. 2008). This high-contract and high spatial resolution data from the NST (diffraction limit of 77 km, pixel scale of 0$''$.0375) allowed us to perform a statistical study of the granule size distribution and explore characteristic scales of solar granulation in detail. The results formulated below do not depend on the thresholding technique: The traditional single-threshold technique and the MLT technique produced similar results.  Analyzing 36 independent granulation images, we concluded that there are two populations of granular structures. First population is comprised of relatively large (regular) granules of the mean size of approximately 1000~km across. The second  a population includes \"mini-granules\" - a continuous power-law population of convective structures that dominate on scales less than approximately 600~km.   The size histogram of regular granules is approximated by a normal (Gaussian) distribution function with a mode of 1050~km and the standard deviation of 480~km. Granules with sizes of 1080-1320~km across (with a mean value of about 1200~km) contribute substantially to the total area of granules (see Figure 2). These regular granules appear to be randomly distributed over the solar surface. We further confirm existence of the characteristic (or ``dominant'' per Roudier \\& Muller, 1986) scale of granules. We found it to be about 1080-1300~km (1$''$.49 - 1$''$.79) from the area contribution function (Figure 2) and 1050$\\pm$480~km from the PDF approximation (Figure 3).   The size distribution of mini-granules can be approximated with a power-law function with a slope of -1.82$\\pm$0.12. Their contribution to the total granular area is comparable to that of regular granules. The mini-granules are mainly confined to broad inter-granular lanes and form chains and clusters intermittent with voids.  The flatness function (a measure of intermittency and multi-fractality) calculated directly (without any thresholding) from the images indicates that solar granulation is non-intermittent on scales exceeding 600 km, and it becomes highly intermittent and multi-fractal on smaller scales. Thus, a random Gaussian-like distribution of granules over the solar surface holds down to 600 km only. On smaller scales, the multi-fractal spatial organization of the mini-granular structures takes over.   Using Hinode data, Yu et al. (2011) reported two types of granules: small and large ones. The dividing point between them is 1044~km, which exceeds significantly that found in our study (600~km). The 1044~km scale is very close to the mode (1050~km) of the Gaussian distribution of the regular granules studied here. More importantly, the minimum size of granules studied in Yu et al. (2011) is half of an arc-second (360~km). In other words, the small granules reported by these authors correspond mainly to the plateau in the size distribution reported here (see Fig. 3), while the subset of mini-granules was mostly missed in their study.  A possible interpretation of mini-granular structures is that they are fragments of regular granules, which are subject to highly turbulent plasma flows in the intergranular lanes, where the intensity of turbulence is enhanced (Nordlund et al. 2009). The association between the mini-granulation and magnetic and velocity fields, as well as efforts to detect mini-granulation in numerical simulations of solar magneto-convection are subjects for future research.  As for now, it is evident that the complex picture of solar near-surface magneto-convection might be even more complex.  We are thankful to anonymous referees whose comments helped us to improve the paper. Authors gratefully acknowledge help of the NST team and support of NSF (ATM-0716512 and ATM-0847126), NASA (NNX08AJ20G, NNX08AQ89G, NNX08BA22G), AFOSR (FA9550-12-1-0066)."
    },
    "1208/1208.0334_arXiv.txt": {
        "abstract": "{We investigate unbound dark matter particles in halos by tracing particle trajectories in a simulation run to the far future ($a=100$).  We find that the traditional sum of kinetic and potential energies is a very poor predictor of which dark matter particles will eventually become unbound from halos.  We also study the mass fraction of unbound particles, which increases strongly towards the edges of halos, and decreases significantly at higher redshifts.  We discuss implications for dark matter detection experiments, precision calibrations of the halo mass function, the use of baryon fractions to constrain dark energy, and searches for intergalactic supernovae.} ",
        "introduction": "In an expanding cosmology, a particle's trajectory depends not only on its initial velocity, but also on the dynamical evolution of nearby structure.  The meaning of ``unbound'' in this case is both nontrivial and physically interesting.  Traditionally, this word describes particles which can escape to arbitrarily large distances from their initial host dark matter halo.  We investigate the basic demographics of unbound dark matter by tracing particle trajectories in cosmological simulations.  In addition, we also investigate traditional unbound particle classification techniques based on kinetic energies.  While intellectual curiosity motivated this work, there are many practical implications.  For example, we find that the fractional abundance of high-energy dark matter particles depends strongly on halo environment; e.g., the distance to the nearest larger halo.  This suggests that the velocity distribution of dark matter depends on halo environment, which is important for dark matter direct detection experiments \\citep{XENON100,CRESST,DAMA,COGENT,CDMS,Mao12,Lisanti11,PS1,PS2,PS3,PS4,PS5,PS6,PS7}.  We therefore test the expected effect from Andromeda on the Milky Way's dark matter velocity distribution directly.  Dark matter particle trajectories also have an effect on the baryon fraction in halos.  When a high-energy dark matter particle leaves the virial radius, the virial mass of the halo is lowered.  The hot halo gas will expand adiabatically in the reduced potential, nominally leaving the hot gas fraction unchanged.  However, the centrally-concentrated stars will not be able to expand, meaning that the overall baryon fraction will increase.  This effect will result in a redshift- and halo-mass-dependent baryon fraction.  Of course, the true picture is made more complicated by a number of effects, including angular momentum exchange, gas cooling, and galaxy feedback.  Nonetheless, it is possible to make a straightforward estimate of its magnitude. In the future, constraining this effect may be important for surveys which aim to place precision constraints on cosmology (e.g., BOSS, DES, BigBOSS, Pan-STARRS, eRosita, Herschel, Planck, JWST, and LSST; \\cite{BOSS,DES,BigBOSS,Pan-STARRS,eROSITA,Herschel,Planck,JWST,LSST}).  In order to fully realize their statistical power, these surveys depend on precision calibrations of the halo mass function to 1\\% or better accuracy \\cite{Wu10,Cunha10}.  Observations of intergalactic supernovae \\citep{Sharon10,Sand11,Barbary12} relate to unbound particles, especially observations of supernovae traveling at high speeds \\citep{GalYam03} and hostless gamma-ray bursts (GRBs) \\cite{Fong13}.  Such escaping objects might either come from the intracluster stellar population \\citep{Monaco06,Barbary12}, from merging events \\citep{Murante07,Teyssier09}, or from hypervelocity stars ejected from the central galaxy \\citep{Napiwotzki11}.  Since stars are effectively collisionless and so evolve dynamically like dark matter particles, one can place upper limits on the fraction of supernova and GRB progenitors ejected through merger events.  In \\S \\ref{s:methods}, we discuss the dark matter simulations we employ as well as the halo finder and conventions used for calculating kinetic and potential energies.  We discuss unique qualitative features of unbound particles in \\S \\ref{s:energy} and derive kinetic escape thresholds; we also test how well kinetic escape thresholds can classify unbound particles in full cosmological simulations.  We present the main results for the population of unbound particles in \\S \\ref{s:results} and discuss how these results impact the science considerations above in \\S \\ref{s:discussion}.  Finally, we summarize our conclusions in \\S \\ref{s:conclusions}.  We use simulations with multiple cosmologies in this work, but our primary results assume a flat, $\\Lambda$CDM cosmology with main parameters $\\Omega_b = 0.04$, $\\Omega_M = 0.27$, $\\Omega_\\Lambda = 0.73$, and $h = 0.7$.  All halo masses at $a\\le1$ are defined using the virial overdensity criterion \\cite{mvir_conv}. ",
        "conclusions": "\\label{s:conclusions}  We have studied the properties of unbound dark matter particles, defined as particles which can travel arbitrarily far from their halo in the far future.  Our main findings are as follows: \\begin{enumerate} \\item On average, at $z=0$, 2--6\\% of particles in halos are unbound (i.e., can escape to infinity; \\S \\ref{s:radial}). \\item The fraction of unbound particles in halos is a very strong function of redshift, and approximately a linear function of distance from the halo center (\\S \\ref{s:radial}). \\item Total kinetic and potential energies cannot be used to predict which particles will become unbound.  Moreover, total kinetic and potential energies are strongly non-conserved at the single-particle level in simulations.  The level of non-conservation is independent of simulation resolution (\\S \\ref{s:full_sims}). \\item The presence of Andromeda does not impact the velocity distribution of dark matter in the Milky Way at the galactocentric radius (\\S \\ref{s:dm_detect}). \\item Halo mergers cannot be used to explain a hostless supernovae or gamma ray burst fraction which exceeds the order of 1\\% (\\S \\ref{s:hostless}). \\item The baryon fraction in massive ($10^{14}$--$10^{15}\\Msun$) halos is negligibly boosted ($0.5-1\\%$) on account of dark matter particles which leave after entering the halo (\\S \\ref{s:baryons}). \\end{enumerate}"
    },
    "1208/1208.0933_arXiv.txt": {
        "abstract": "We revisit the calculation of the Ohmic dissipation in a hot Jupiter presented in \\citet{Laine} by considering more realistic interior structures, stellar obliquity, and the resulting orbital evolution. In this simplified approach, the young hot Jupiter of one Jupiter mass is modelled as a diamagnetic sphere with a finite resistivity, orbiting across tilted stellar magnetic dipole fields in vacuum. Since the induced Ohmic dissipation occurs mostly near the planet's surface, we find that the dissipation is unable to significantly expand the young hot Jupiter. Nevertheless, the planet inside a small co-rotation orbital radius can undergo orbital decay by the dissipation torque and finally overfill its Roche lobe during the T Tauri star phase. The stellar obliquity can evolve significantly if the magnetic dipole is parallel/anti-parallel to the stellar spin. Our results are validated by the general torque-dissipation relation in the presence of the stellar obliquity. We also run the fiducial model in \\citet{Laine} and find that the planet's radius is sustained at a nearly constant value by the Ohmic heating, rather than being thermally expanded to the Roche radius as suggested by the authors. ",
        "introduction": "In the study of the orbital distribution of known Jupiter-mass exoplanets, the radial-velocity method has revealed a pile-up of hot Jupiters with orbital periods of $\\sim$ 3 days (e.g., see The Extrasolar Planets Encyclopedia website at http://exoplanet.eu). A number of models have been proposed to explain the pile-up, such as an inner disk cavity stopping planet migration \\citep[e.g.][]{Lin96,Rice,Benitez}, and tidal circularization of a gas giant planet in an extremely eccentric orbit arising from planet-planet interactions after the proto-planetary disk disperses \\citep[e.g.][]{Chatterjee,Nagasawa,WL11,Naoz}. In addition, tidal heating in a young hot Jupiter in a moderately eccentric orbit may inflate the planet over its Roche-lobe, resulting in mass loss and therefore leading to the halting of planet migration or even planet destruction \\citep[e.g.][]{Gu03,Gu04,Chang}. However the excitation of the planet's eccentricity in this case is subject to the uncertain density profile of the inner edge of  a disk \\citep{Rice,Benitez}.  Aside from the tidal dissipation that relies on the presence of orbital eccentricity, \\citet{Laine} invoked the Ohmic dissipation to inflate a young hot Jupiter in a circular orbit by adopting the model proposed by Campbell (1983,1997) for the magnetic interactions in the AM Herculis systems. In this simplified model, the planet is assumed to behave as an imperfect conductor without its own ionosphere and magnetosphere; namely, a diamagnetic sphere with a finite resistivity. In addition, it is assumed that the stellar spin is aligned with the planet's orbit. Since the vacuum space is assumed between the star and the planet, the stellar magnetic dipole must be misaligned to induce electric currents and magnetic fields as the planet circles its T Tauri star. The magnetic torque arises from the Ohmic dissipation in the planet at the expense of the spin-orbit energy (see the \\S2.3).  Normally a planet even without its own fields possesses an ionosphere due to the photo-ionization of the upper atmosphere. An induced magnetosphere can form above the ionosphere \\citep{Zhang}. As a planet orbits a star with a tilted magnetic dipole, the ionosphere may shield the time-varying stellar fields so sufficiently that little electromagnetic field can be induced in the planet's interior by the external stellar fields. Nevertheless, one may argue that a young hot Jupiter is already tidally locked by its parent star such that most of its permanent night-side lacks an ionosphere. This argument neglects the global circulation in the atmosphere \\citep[e.g.][]{Showman08,Showman09,TC10,RM10,Dobb,Perna}, which may maintain an ionosphere on the permanent night side. Based on the radio-sounding results from the Venus Express spacecraft, the ionosphere on the night side of Venus is weaker and possibly more sporadic than that on the day side \\citep{Patzold}. Hereafter, we boldly apply the model for AM Her binaries to the entire planet and also consider a vacuum space outside the planet and the star to simplify the calculation. The consequence of this simplification is that other electromagnetic effects such as unipolar induction \\citep{GL69,LL12}, Alfven-wave wings \\citep{Neubauer,Kopp}, dynamical friction \\citep{Papa}, stellar winds \\citep{Vidotto}, planetary winds \\citep{Adams,Trammell}, stellar fields diffusing into the planet interior \\citep{Campbell05}, and magnetic reconnections are all ignored (also see Lanza 2011 for a recent review). In addition, any electromagnetic effects associated with atmospheric circulations are not being considered for further simplicity \\citep{Perna,BS10,BSB,WL12}. In short, we restrict ourselves only to the diamagnetic part of the star-planet magnetic interaction\\footnote{The same concept has been applied to star-disk magnetic interactions \\citep{Lai99} in which the magnetic response of the disk is modelled by a diamagnetic disk as well as a magnetically threaded disk. Unlike our model planet possessing a finite resistivity, the disk is assumed to be a perfect conductor in the diamagnetic part of their model.}, as was modelled by \\citet{Laine}. It should be noted that the Ohmic heating proposed by \\citet{Laine} is short-lived, since the stellar magnetic fields decay significantly during the T Tauri phase. This is in contrast with the Ohmic dissipation model by \\citet{BSB}, which is long-lived. Consequently, this study is concerned exclusively with the early evolution of hot-Jupiter systems.  The observations using the Rossiter-McLaughlin effect \\citep{Ohta} suggest that dwarf stars hosting transiting planets may have possessed a wide range of stellar obliquities \\citep[e.g.][]{Winn10,Winn11}. These observational findings seem against the conventional paradigm in which a planet should orbit in the same direction as the stellar spin as the star and planets form together in a proto-planetary disk. A number of N-body numerical simulations demonstrated that after the proto-planetary disk disperses, planet-planet interactions accompanied by tidal circularization, as mentioned in the first paragraph of the Introduction, can generate obliquities. It was also proposed that before the proto-planetary disk dissipates, the warp torque resulting from the magnetic interactions between the proto-star and the inner part of the disk would move the stellar spin away from the disk angular momentum despite the presence of gas accretion onto the proto-star \\citep{Lai11,FL11}. Motivated by the latter works, it is timely to consider a more complex case in which stellar obliquity $\\lambda$ is not zero; i.e., the orbital axis is not aligned with the stellar spin.  It should be noted that the tidal dissipation in the star drives the system to the spin-orbit alignment as well as synchronization \\citep[e.g.][]{Hut,Soko,Lai11}. To make the problem tractable, we do not take account of the influence on $\\lambda$ driven by the proto-planetary disk or by tidal interactions with the proto-star, but simply take $\\lambda$ as a free parameter in this work. In addition, we assume that the planet spin is tightly being synchronized with its orbital motion during the evolution, therefore generating negligible dissipation in the planet \\citep[e.g.][]{Gu03}. This simplification allows us to ignore the effect due to the planet spin in the calculation.  Owing to the Ohmic dissipation and the resulting magnetic torques, the stellar spin, planet's orbit, and the interior structure of the planet evolve simultaneously. To calculate the coupled evolution more precisely, we adopt an interior-structure model \\citep[for details see][]{Chang} to compute the planet resistivity and the thermal response of the planet due to the Ohmic heating.  The structure of the paper is organized as follows. In \\S2, we describe the equations for the coupled secular evolutions of stellar spin, planet's orbit, and planet interior structure due to the diamagnetic interaction between a young hot Jupiter and its parent T Tauri star. In order to understand the dependence of Ohmic dissipation on various orientations of the stellar spin and magnetic dipole moment, we first conduct a parameter study in \\S3 to investigate this with no secular evolutions. The parameter study involving secular evolutions is then presented in \\S4. Finally, we summarize and discuss the results in \\S5. ",
        "conclusions": "We revisit the magnetic interaction between a hot Jupiter and its T Tauri star investigated by \\citet{Laine}. In the original work, the authors considered a Roche-lobe sized hot Jupiter without its own magnetic fields. The stellar spin was assumed to be aligned with the orbital axis; i.e. the stellar obliquity $\\lambda=0$. As the planet orbits its parent star, the Ohmic dissipation in the planet is induced by the stellar magnetic dipole tilted away from the stellar spin with an angle $\\alpha$. To calculate the electric resistivity of the planet, the interior structure was modelled as a sphere consisting of a polytropic core and an isothermal outer layer. In their model, the planet lies outside the co-rotating orbital radius such that the forcing frequency $\\omega_->0$. Based on their fiducial model, the authors suggested that the dissipation torques are not able to cause any significant orbital change. Nevertheless, the Ohmic heating occurring in the outer part of the planet is intense enough to inflate the planet up to the Roche radius. The mass loss through the Lagrangian 1 point toward the central star provides the angular momentum to the planet and thus possibly halts the planet migration in the disk.  Motivated by a wide range of the stellar obliquity detected in hot-Jupiter systems \\citep[e.g.][]{Winn11}, which in theory could be excited during the T Tauri phase \\citep{Lai11,FL11}, we extend the original model by considering the coupled evolution of the interior structure, planet's orbit, and the stellar spin in the presence of the stellar obliquity $\\lambda$. We focus on the secular evolution due to the dissipation torques and show that the Ohmic dissipation in the planet can be contributed linearly from the forcing associated with 4 frequencies: $\\omega_+$, $\\omega_*$, $n$, and $\\omega_-$. Owing to the complication of the problem involving multiple frequencies, we begin with a couple of test runs based on a given interior structure for the dissipation calculation, which are further validated by the general torque-dissipation relation given by Equation~(\\ref{eq:heat_torque}) as well as the skin-depth estimation using Equation~(\\ref{eq:skin_depth}).  The coupled evolutions are then carried out for a number of cases listed in Table~\\ref{tbl-1} for the purpose of parameter studies. The evolutions are computed from $t=0.7$ to $10$ Myrs. The radius of the coreless hot Jupiter of $1M_J$ is $2.045R_J$ at the beginning. Initially, the T Tauri star is assumed to spin at a rate $\\omega_*=n/2$ to imitate the final stage of the planet migration scenario with a giant planet inside the magnetospheric cavity of the disk (Lin et al. 1996; Rice et al. 2008; cf. Ben\\'itez-Llambay et al 2011). Since the planet lies inside the co-rotating orbit, the planet continues to migrate inwards due to the Ohmic dissipation in the planet. Without modelling the star-disk magnetic interactions, the co-rotating orbit is simply assumed to correspond to the inner edge of the magnetically truncated disk despite the presence of $\\lambda$ and $\\alpha$. This is certainly one of the limitations of the study. In most of the cases, $\\omega_*$ is assumed constant to simply resemble any processes, such as disk locking, that maintain the stellar spin.  Three initial orbital distances $D_i \\approx 0.02$, 0.03, and 0.04 AU are considered. With our input parameters for $D_i\\approx 0.02$ AU, the intense dissipation confined near the planet surface only enlarges $R_p$ by $< 3$\\%. Nonetheless, the dissipation torques decay the orbit to the Roche zone in just a few $10^5$ to about 1 million years. The torques also evolve $\\lambda$ when $\\lambda_i \\neq 0$. Since $M_pD^2 n \\ll I_* \\omega_*$, $\\dot \\lambda$ is primarily contributed from the movement of the orbital axis rather than the stellar spin axis. When $\\alpha$ is zero, there exists an unstable equilibrium for the orientation of the stellar spin axis, which points roughly about $90^\\circ$ from the orbital axis. Consequently, the dissipation torques direct the orbital axis toward the stellar spin for a prograde orbit but toward the anti-parallel direction to the spin for a retrograde orbit. When $\\alpha$ is non-zero, the orbital axis and the stellar spin can either evolve toward alignment/anti-alignment for small $\\alpha$ or become more misaligned for $\\alpha \\lesssim 90^\\circ$. Because the stellar spin is not spun up/down by the dissipation torque when $\\alpha=0$, it follows that the stellar obliquity evolves more quickly when the stellar spin is parallel/anti-parallel to the stellar magnetic dipole.  For the young hot Jupiter initially at the farther distance $D_i \\approx $ 0.03 AU, the dissipation is modest but still strong enough to more or less sustain the initial $R_p$ except for the cases with $\\alpha=0$. The relatively fast decrease of $\\lambda$ for $\\alpha=0$ weakens the dissipation and the resulting torques, leading to the planet contraction and slow orbital decay in the late stage of the evolution. Therefore, in terms of the cases we have studied, the planet with $\\alpha \\neq 0$ undergoes substantial orbital decay in a few Myrs and finally overflows the Roche-lobe, while the planet with $\\alpha=0$ can shrink its orbit but not sufficiently to allow for Roche-lobe overflow.  Owing to the weaker interaction at the larger orbital distance, the planet of 1 $M_J$ in all the cases starting from $D_i \\approx 0.04$ AU contracts and then roughly reaches quasi-thermal equilibrium during the T Tauri phase, with the final $R_p$ smaller than those in the cases for $D_i\\approx$ 0.02 and 0.03 AU. The corresponding orbital decays and the changes in the stellar obliquity are substantially smaller. The planet moves barely from its initial orbit and thus is not able to reach its Roche-lobe. We also carry out the simulation for the fiducial model in \\citet{Laine} and find that the Ohmic heating can only sustain the radius of the less massive young hot Jupiter ($M_p=0.63 M_J$), rather than thermally expanding the planet to its Roche radius as suggested by the authors.   The induced dissipation rates are as high as $10^{30-31}$ erg/s when the planet moves to about $D <$ 0.02 AU. The intense heating near the planet's surface does not generate local temperature inversion in our model, in contrast to the surface heating models presented in \\citet{Gu04} and \\citet{WL12}. It is probably because the dissipation responsible for the temperature inversion lies within a thin shell fairly deep in Gu et al. 2004 (a prescribed narrow Gaussian region) and in Wu \\& Lithwick 2012 (a region at about the optical depth of 100), while the dissipation in our diamagnetic induction model is deposited so close to the surface that it is easily lost and so there is no local maximum in T.  In this work, we introduce the planet at $t_i=0.7$ Myrs and adopt a constant magnetic dipole moment. In theory, a gas giant planet can form later and migrate to the magnetospheric cavity at a later time \\citep[e.g.,][]{il09,Mordasini}. Besides, the magnetic dipole moment may decay over the course of a few Myrs \\citep{JK,YJ11}. Hence, our results probably give the suggestive values for the maximum changes of the stellar obliquity, orbital distance, and planet radius during the T Tauri star phase.  The orbital decay of a young hot Jupiter in our magnetic model is not significant unless the distance to the T Tauri star is smaller than about 0.03 AU. Some other process, such as gravitational tides \\citep{Chang} or a small magnetospheric cavity during FU-Orionis outbursts \\citep{Baraffe09,Adams09}, can bring the planet in to that distance. The Ohmic mechanism could be the final stage in bringing a planet in very close to the T Tauri star or even leading to a Roche-lobe overflow. This may provide one of the explanations for the pile-up of hot Jupiters with the orbital periods of $\\sim$ 3 days, and could also reduce the too-high population of hot Jupiters inferred from the population synthesis model \\citep{il09}. In this work, we do not study the post-evolution of a Roche-lobe filled planet. In terms of our model parameters, a young hot Jupiter located within 0.03 AU from its T Tauri star can undergo fast orbital decay on the timescales much shorter than 10 million years. Therefore it is possible that the young planet overflows its Roche lobe, migrates out, then migrates in, and overflows again. Consequently, the planet suffers from intermittent mass losses until its density is low enough to go through the stage of the runaway adiabatic mass loss \\citep{Chang}, leading to the demise of the planet during the T Tauri phase.   As has been described in the Introduction, there is a large body of literature devoted to a variety of magnetic interactions between a hot Jupiter and its parent star, some of which can also cause the angular momentum to transfer between the planet's orbit and the stellar spin. The efficiency of angular momentum transfer is model dependent, relying on the Ohmic dissipation rate. In the T Tauri phase, the presence of a disk is expected to magnetically affect the stellar spin and perhaps the planet's orbit. \\citet{Lai11} and \\citet{FL11} considered a hybrid magnetic model including diamagnetic induction and magnetic-field linkage for the purpose of the generation of the stellar obliquity. In the study presented here, we follow the work by \\citet{Laine} and therefore focus only on the diamagnetic interaction between the planet and its T Tauri star. Our simple model suggests that the stellar obliquity starting from a non-zero value may further evolve after the planet migrates into the magnetospheric cavity, making the orbit of the young hot Jupiter incline with the disk plane. As a result, a hot Jupiter does not necessarily lie on the same orbital plane with the planets farther out from the central star. Whether or not our model can provide a wide range of stellar obliquities at the end of T Tauri phase depends on the initial distribution of stellar obliquity as well as the distribution of the direction of stellar dipole moment relative to the stellar spin.  In the Introduction, we also caution that the skin depth beneath the photosphere is one of the major uncertainties of the model. In the presence of a planetary ionosphere or magnetosphere, the value may be appreciably smaller than what we compute in this work. Nevertheless, as an analog of the star-disk magnetic interactions \\citep{Lai99}, our diamagnetic model and other magnetic interactions should be considered together for the orbital evolution inside the magnetospheric cavity. While theoretical models are under development, it is conceivable in the future that photometric variability on timescales of a few days \\citep[e.g.][]{Bouvier} and spectropolarimetry applied to T Tauri stars \\citep[e.g.][]{Donati,Long} would serve as possible detection methods, to search for the variability modes and magnetic perturbations that are associated with the orbital motion of such a young hot Jupiter during the T Tauri star stage."
    },
    "1208/1208.0791_arXiv.txt": {
        "abstract": "% The claim of a neutrino velocity different from the speed of the light, made in September 2001 by the Opera experiment,  suggested the study of the time delays between TeV underground muons in the Gran Sasso laboratory using the old data of the MACRO experiment, ended in 2000. This study can give also hints on new physics in the particle cascade produced by the interaction of a cosmic ray with the atmosphere. ",
        "introduction": "In September 2011 there was a measurement  of the  speed of neutrino faster than the speed of light  by $\\frac{v-c}{c} = 2.48 \\pm 0.28(stat)\\pm 0.30(sys) \\times 10^{-5}$ (Adam, 2011). After many checks we know now that this result was due to hardware problems and  the Opera 2012 result is that the speed of the neutrinos traveling from CERN to the Gran Sasso is $\\frac{v-c}{c} = -0.7 \\pm 0.5(stat) ^{+2.5}_{-1.5}(sys) \\times 10^{-6}$ (Dracos,  2012). This result is in agreement with the results of the other Gran Sasso experiments (Bertolucci,  2012).  However the interest for this claim suggested the possibility to compare neutrino and muon velocity in a cosmic ray cascade. (Montaruli and Ronga, 2011). The interaction of a primary cosmic ray  with the atmosphere produces a cascade with many kind of particles, and in particular neutrinos and muons. Muon neutrinos and muons are produced  mainly via the decay of charged pions and kaons produced in the primary cosmic ray interactions. Above about 10 TeV they can come also from prompt decays of charmed hadrons. This  component has not yet been observed. In a deep underground detector only muon and neutrino are detected. If the neutrino velocity is different from $c$ the  neutrinos in this cascade,  should arrive with times  different  from the times of the muons from the same parent decay, or from another decay, with a time delay that should change according to the neutrino path length that depends on its zenith angle $\\theta$. In underground detectors muon neutrinos are detected looking for induced muons produced by neutrino charged current interactions in the rock, or in  the ice around or inside the instrumented region. Hence, a time spread should be observed between the muons produced  directly by the pion or kaon decay  and the muon produced by neutrino interactions.  The  path length from the meson decay point is a few tens of kilometers for vertical neutrinos  and up to $\\sim 300$~km for near horizontal neutrinos. Assuming the original  time difference  observed in OPERA  nearly horizontal neutrinos should arrive up to 28 nsec before the other secondaries. In  (Gaisser  and Stanev,1998) a table of average production heights neutrinos in the atmosphere has been reported. The typical production height for neutrinos of energy above 20 GeV can be 17.6 km at the vertical, 94.9 km at $\\cos\\theta = 0.25$ and 335.7 km at $\\cos\\theta = 0.05$, which would correspond to 1.4, 78 and 27.6~ns.   There are already limits of tachions or anomalous delayed particles in cosmic rays. The limits are  obtained searching for example signals before  or after the main front of the electromagnetic shower. But this kind of searches stopped some time ago and the last particle data book  review of those data is the one of 1994 (Montanet  et al, 1994). The  limits obtained are of small interest in the framework of the OPERA result. However, if neutrinos were tachions, it is likely that other kind of tachions could exist and this search in very high energy cosmic rays could have a new interest.  It is important to note that the Gran Sasso mountain minimum depth $\\sim 2700~gr/cm^2$ correspond to a minimum muon energy of  1.4 TeV. It easy to compute that, requiring a minimum threshold of 50 MeV in the detector the time difference between two muons underground should be $\\ll ~0.2~nsec$. Therefore anomalous time differences should be a signal of \"new physics\", for example  signal of a supersymmetric massive particles produced in a cosmic ray cascade. For example, let we assume an hypothetical hadron of mass 100 GeV, produced by a an interaction of a proton with center of mass energy 7 TeV (the LHC energy). If this hypothetical hadron interacts or decays  after 10 Km producing at the end muons, the delay between mountain muon from the massive particle and the muon produced in the primary vertex is of the order of 13 nsec. LHC experiments have put limits for new  hadron-like massive particles (Chatrchyan, 2012 and  Aad,  2012)  but it is important to remember that the cosmic ray energy could be larger than the LHC energy. Under Gran Sasso the fraction of multiple muons produced by cosmic ray with center of mass energy $\\ge7~TeV$ is estimated of the order of $10^{-3}$ in MACRO, corresponding to several thousand multiple muon  events in the MACRO data set. One should also consider the possibility that new massive relic particles are directly in the primary cosmic radiations.  The MACRO experiment has done several searches for possible anomalies of the time differences between  muons (Ahlen et al, 1991). The search  was done mainly to study time differences of the order of a few msec or more, but this paper contains also the study of time differences at the ns-level. The statistics was limited to 35832 tracks in events with two or more tracks. In 1992 none was thinking to  tachionic neutrinos and therefore there was no estimate of the number of tracks due to down-going neutrino together with  a primary muons. In (Scapparone,1995) this study was extended to  about 140000 tracks of multi muon events, corresponding to about $4\\%$ of the total MACRO statistics. The time distribution was in agreement with the predictions.  \\begin{myfigure} \\centerline{\\resizebox{130mm}{!}{\\includegraphics{Ronga_2012_1_fig1.eps}}} \\caption{Event with six parallel muons  in 3 MACRO \"supermodules\". On the top there is the full MACRO display, on the bottom the zoom of the 3 supermodules interested by the event. The 12 steamer tube horizontal planes are shown as horizontal lines, the back points are the streamer tubes fired; the scintillator boxes fired are shown as rectangles.} \\label{author-fig1} \\end{myfigure}  In this paper I present an analysis of the full MACRO statistics. This was not an easy job. The main reason is that MACRO ended in 2000 and that most of the analysis software was designed for VAX/AlphaVAX computers and data formats around 1990 (a geological era for computers!). A lot of time has been necessary to convert programs and to find data files, some time stored on data tape cassette of old formats, obsolete and not supported by modern computers. ",
        "conclusions": "This work ended some time after the solution of the superluminal neutrino puzzle, however I think that has been very useful to remember that cosmic rays are still important tools in particle physics. For the superluminal neutrino 2 tracks were expected 0 were found.  Considering the different mean-lifes of the pion and the kaon, an \"exotic\" limit can be derived from  the horizontal tracks of Fig 2 on the equality of the pion and kaon speed in a cascade produced by a primary with $E\\ge 3TeV$: $|\\beta_{\\pi}-\\beta_{k}| \\ll 1.5 \\times 10^{-4}$. This result is at the moment of very low interest but the superluminal neutrino saga has shown that nothing could be given as guaranteed. More investigations are necessary on the  delayed tracks in events with multiplicity bigger than 3  and on  massive particles in cosmic rays.  This work has shown once again the importance to save past experiment data for further analysis.  I must thank the  MACRO collaboration that built and run the detector and many MACRO peoples that helped me to recover data and programs and particularly Nazareno Taborgna of the Gran Sasso laboratory, that has been able to save a working alphaVAX with several MACRO original disks. A particular thank is to Teresa Montaruli for useful and deep discussions."
    },
    "1208/1208.1873_arXiv.txt": {
        "abstract": "In the last few years  many globular clusters (GCs) have revealed complex color-magnitude diagrams, with the presence of multiple main sequences (MSs), broaden or multiple  sub-giant branches (SGBs) and MS turn offs, and broad or split red giant branches (RGBs).  After a  careful correction for differential reddening, high accuracy photometry with the  Hubble Space Telescope presented in this paper reveals  a  broadened or even split SGB in five additional Milky  Way  GCs: NGC 362, NGC 5286,  NGC 6656,  NGC 6715, and  NGC 7089.  In addition, we confirm (with new and archival HST data) the presence of a split SGB in 47Tuc, NGC 1851, and NGC 6388.  The fraction of faint SGB stars with respect to the entire SGB population varies from one cluster to another and ranges from $\\sim$0.03 for NGC 362 to $\\sim$0.50 for NGC 6715.  The average magnitude difference between the bright SGB and the faint SGB is almost the same at different wavelengths.  This peculiarity is consistent with the presence of two groups of stars with either an  age difference  of about 1-2 Gyrs,  or a significant difference in their overall C+N+O content. ",
        "introduction": "\\label{introduction} Recent photometric studies have revealed unexpectedly complex color-magnitude diagrams (CMDs) in many Globular Clusters (GCs), indicating that these stellar systems are not as simple as we have been assuming for decades (see Piotto 2009 for a recent review).  $\\omega$ Centauri is the most famous example of a GC hosting complex multiple stellar populations. It has been known since  the late nineties that its CMD shows  multiple red-giant branches (RGBs, Lee et al.\\ 1999, Pancino et al.\\ 2000), sub-giant branches (SGBs e. g.\\  Bellini et al.\\ 2010) and  multiple main sequences (MSs, Anderson et al.\\ 1997, Bedin et al.\\ 2004).  To date, multiple or broad RGBs have been observed in nearly all the GCs that have been observed with good signal-to-noise in the appropriate photometric bands (e. g.\\ Marino et al.\\ 2008, Yong et al.\\ 2008, Lee et al.\\ 2009, Lind et al.\\ 2011).  Recent studies also suggest that multimodal MSs could be quite common among GCs.  In addition to the spectacular case of NGC 2808, which shows three  distinct MSs (Piotto et al.\\ 2007), double and triple  MSs have been observed in several GCs, including NGC 104, NGC 6752, and NGC 6397 (Anderson et al.\\ 2009, Milone et al.\\ 2010, 2012a,b) and have been associated to stellar generations with a different content of helium and light-elements (e. g.\\ Norris 2004, Piotto et al.\\ 2005, D'Antona et al.\\ 2005).  Also the CMDs of many stellar clusters in the Large and the Small Magellanic Cloud (LMC,  SMC) are not  consistent with single stellar populations (Bertelli et al.\\  2003, Mackey  et al.\\ 2008).  Milone et al.\\  (2009a) analyzed sixteen intermediate age clusters in the LMC and found that at least eleven of them (i.e.\\ about the 70 \\% of the whole sample) exhibit a  wide spread or a split around the main sequence turn-off (MSTO), which is consistent with the presence of multiple or prolonged star formation episodes.  Spectroscopic studies have long shown that most of the GCs exhibit some correlations and anticorrelations among light-element abundances (such as the Na-O anticorrelation, Kraft et al.\\ 1979, 1994, Ramirez \\& Cohen 2002). The fact that light-element variations have also been observed among unevolved stars (e.g.\\ Cannon et al.\\ 1998, Gratton et al.\\ 2001) indicates that they have a primordial origin (see Gratton et al.\\ 2004 for a review).  However, it must be noted that spectroscopic analysis is necessarily limited to a small sample of (bright) stars, and therefore allows us to identify stellar generations which may constitute only a small fraction of the cluster populations. Often spectroscopy investigations cannot follow multiple stellar generations in all evolving sequences (in particular along the main sequence and the white dwarf cooling sequence). Information on multiple stellar generations is maximized when spectroscopic and photometric data can be used together. Otherwise, we  must rely on photometry for a complete analysis of multiple populations (relative fraction, radial distribution, main chemical properties), which manifest themselves in different ways, in different clusters and different evolutionary parts of the CMD.  The multiple populations in NGC 1851 and NGC 104 manifest themselves most prominently in terms of a splitting of the SGB (Milone et al.\\  2008, 2012a, Anderson  et al.\\ 2009).  NGC 6388 also shows hints of such a splitting as well (Piotto 2008 and Moretti et al.\\ 2009).  These SGB splits have been interpreted as indicating two groups of stars with either an age difference  of about  1-2 Gyrs, or a significant difference in their overall C+N+O content (Cassisi et al.\\ 2008, Ventura et al.\\ 2009, Di Criscienzo et al.\\ 2010).  In this paper we present detection of the broad (and most likely split) SGBs in five additional  Milky Way GCs: NGC 362, NGC 5286,  NGC 6656, NGC 6715, and  NGC 7089, and confirm the presence of a double SGB in NGC 6388, NGC 104, and NGC 1851. In a companion paper based on the spectroscopy of stars selected from the CMDs published in this work, we have found that the two SGBs of NGC 6656 (M22) are made of stars with a different content of iron, $s$-process elements, and C+N+O (Marino et al.\\ 2012). A bimodality in $s$-elements has been also observed along the RGB of NGC 1851 and associated with the double SGB of this cluster (e.g.\\ Yong et al.\\ 2008, Villanova et al.\\ 2010, Lardo et al.\\ 2012).  This paper is organized as follows.  In Sect.~\\ref{data} we describe the data set, the photometric reduction, the zero point calibrations, the selection of stars with high-quality photometry, and the differential-reddening correction.  The initial detection of the  SGB-broadenings is presented  in Sect.~3.  We then proceed to study the SGBs in increasing detail.  We first use new and archival material: $i)$ to derive proper motions to exclude field objects (Sect.~4) and $ii)$ to provide confirmation of the SGB-broadenings from several independent data-sets (Sect.~5).  Section~6 then examines these SGBs through a variety of bandpasses.  Finally, in Sect.~7, for each cluster we compare the fractions of stars in the different SGBs with the distribution of stars along the HB in hopes of identifying the same populations in the different evolutionary branches (Sect.~8).  Section~9 contains the summary of results and an examination of open questions. ",
        "conclusions": "In this paper we have shown that the Galactic GCs NGC 362, NGC 5286, NGC 6388, NGC 6656, NGC 6715, and NGC 7089, exhibit double or broadened SGBs, similar to those identified in NGC 1851 (Milone et al.\\ 2008) and NGC 104 (Anderson et al.\\ 2009, Milone et al.\\ 2012a).  When we compare different CMDs of the same cluster made with different magnitudes and color combinations, we find that the magnitude difference between the bright and the faint SGB components remains approximately constant and does not depend on the used filters. Therefore, we conclude that the split/spread SGB of the GCs studied in this paper can be interpreted in terms of two stellar groups with either a difference in age by $\\sim$1-2 Gyr or a large  difference in the total C+N+O abundance, as previously suggested by Cassisi et al.\\ (2008) and Ventura et al.\\ (2009) for the case of NGC 1851.  In the cases of NGC 362, NGC 7089, and NGC 5286, the small number of fSGB stars do not allow us to firmly estabilish if the two groups of fSGB and bSGB stars correspond to two distinct stellar populations or if the poorly populated fSGBs are just the tail of an exthended star-formation history.  The fractions of faint and bright SGB stars with respect to the total number of SGB stars changes from cluster to cluster, and ranges from (0.03:0.97) in the case of NGC 362 to (0.51:0.49) for the Sagittarius dwarf galaxy's GC NGC 6715."
    },
    "1208/1208.6512_arXiv.txt": {
        "abstract": "We reproduce the mid-infrared to radio galaxy counts with a new empirical model based on our current understanding of the evolution of main-sequence (MS) and starburst (SB) galaxies. We rely on a simple Spectral Energy Distribution (SED) library based on \\textit{Herschel} observations: a single SED for the MS and another one for SB, getting warmer with redshift. Our model is able to reproduce recent measurements of galaxy counts performed with \\textit{Herschel}, including counts per redshift slice. This agreement demonstrates the power of our 2 Star-Formation Modes (2SFM) decomposition for describing the statistical properties of infrared sources and their evolution with cosmic time. We discuss the relative contribution of MS and SB galaxies to the number counts at various wavelengths and flux densities. We also show that MS galaxies are responsible for a bump in the 1.4~GHz radio counts around 50~$\\mu$Jy. Material of the model (predictions, SED library, mock catalogs...) is available online\\footnote{at http://irfu-i.cea.fr/Pisp/matthieu.bethermin/}. ",
        "introduction": "Recent observational studies have shown that two distinct star-forming (SF) mechanisms are required to describe the SF galaxy population. The so-called SF main sequence (MS) is composed of secularly-evolving galaxies that display a tight correlation between stellar mass ($M_\\star$) and star formation rate (SFR) at a given redshift \\citep[e.g.][]{Elbaz2007,Noeske2007,Daddi2007}. This population accounts for $\\sim$85\\% of the star formation rate density (SFRD) in the Universe \\citep{Rodighiero2011,Sargent2012} at z$<$2. The rest of the star-formation budget is provided by starbursts (SB), i.e. galaxies with very high specific star formation rates (sSFR=SFR/$M_\\star$), probably induced by recent mergers \\citep[e.g.][]{Elbaz2011,Rodighiero2011}. Recently, \\citet[][S12 hereafter]{Sargent2012} showed that infrared (IR) luminosity functions (LF) can be reproduced by jointly considering the mass function of SF galaxies (SFMF), the evolution of the sSFR of MS galaxies, and its distribution at fixed M$_\\star$, with a separate contribution from MS and SB galaxies.\\\\  Wavelength-dependent galaxy number counts are an additional, important constraint for evolutionary models of infrared galaxies. While purely semi-analytical models \\citep[e.g.][]{Lacey2010,Somerville2011} struggle to reproduce infrared (IR) number counts, phenomenological or hybrid models \\citep[e.g.][]{Bethermin2011,Gruppioni2011,Rahmati2011,Lapi2011} fare better but are in general descriptive and use an evolution of the luminosity function which is not motivated by physical principles. However, these recent models which reproduce the total counts passably, are excluded at $>$3\\,$\\sigma$ by the recent \\textit{Herschel} measurements of counts per redshift slice \\citep{Berta2011,Bethermin2012}. This shows how important redshift-dependent constraints are to accurately model the evolution of galaxies, and motivates the development of a new generation of models.\\\\    We present a new model of IR galaxy counts which builds on the 2-Star-Formation-Mode framework (2SFM) S12 introduced. This fiducial model is intuitive and based on our current observational knowledge of the evolution of MS and SB galaxies. All model parameters are constrained by external datasets and require no additional fine-tuning. We assume a Salpeter initial mass function and a \\textit{WMAP}-7 cosmology.\\\\ ",
        "conclusions": "\\label{sect:conclusion}  Our model based on the main assumption of two SF modes (MS and SB) is able to accurately reproduce the emission of galaxies integrated over most of the Hubble time as probed by galaxy counts from the mid-IR to radio wavelengths. This model contains two main ingredients: the evolution of MS and SB galaxies based on the S12 formalism and a new library of MS and SB SEDs derived from \\textit{Herschel} observations (M12). Despite its simplicity, our model provides one of the best fits achieved so far to the number counts, including counts per redshift slice in the SPIRE bands, which were poorly reproduced by the previous generation of models. All these results were obtained without any arbitrary tuning of parameters that are not constrained by observations, contrary to most previous models. The decomposition into 2 modes of SF (2SFM), i.e. MS and SB, associated with two different families of SEDs, is thus a very powerful framework to statistically describe the dust emission of galaxies across cosmic time. In addition, we present a new stochastic AGN treatment, and also found that MS galaxies are responsible for a bump in the 1.4~GHz radio counts around 50~$\\mu$Jy.\\\\   This model can be combined with halo models assuming a link between SFR, M$_\\star$, and halo mass \\citep[e.g.][]{Bethermin2012a,Wang2012} to interpret the clustering of infrared galaxies and the fluctuation of the cosmic infrared background \\citep[e.g.][]{Planck_CIB}. Finally, this model and its future extensions will provide predictions for the next generation of IR, mm, and radio surveys, and, in particular to anticipate which galaxy populations will be preferentially detected, depending on the survey strategy adopted.\\\\"
    },
    "1208/1208.4517_arXiv.txt": {
        "abstract": "It has recently been shown that a significant fraction of late-type members of nearby, very young associations (age $\\lesssim$10 Myr) display excess emission at mid-IR wavelengths indicative of dusty circumstellar disks. We demonstrate that the detection of mid-IR excess emission can be utilized to identify new nearby, young, late-type stars including two definite new members (``TWA 33\" and ``TWA 34'') of the TW Hydrae Association.  Both new TWA members display mid-IR excess emission in the Wide-field Infrared Survey Explorer (WISE) catalog and they show proper motion and youthful spectroscopic characteristics -- namely H$\\alpha$ emission, strong lithium absorption, and low surface gravity features consistent with known TWA members.  We also detect mid-IR excess -- the first unambiguous evidence of a dusty circumstellar disk -- around a previously identified UV-bright, young, accreting star (2M1337) that is a likely member of the Lower-Centaurus Crux region of the Scorpius Centaurus Complex. ",
        "introduction": "Low-mass stars are the most abundant stellar constituent in our Galaxy and are likely the typical planet hosts, thanks to recent radial velocity and microlensing results (e.g., \\citealt{gau08}).  The diverse zoo of planetary systems discovered to date implies that conditions in circumstellar disks are intimately linked to the final planetary system architectures. In turn, circumstellar environments appear to be controlled by the mass of the central star they orbit (Pascucci et al. 2009 and references therein).  Thus, it is of great interest to understand in detail the evolution of circumstellar material surrounding low-mass stars.  In their analysis of mid-IR excess fractions for young, nearby associations, \\cite{sch12} show that a significant number of M-type stars display mid-IR excess emission in stellar groups younger than $\\sim$10 Myr.  \\cite{riaz12} found a similar result in their analysis of primordial disk fractions for young clusters, namely that disks around later-type stars remain in the primordial stage for a longer period of time than disks around stars of earlier spectral types.  Based on the WISE channel 4 at 22 $\\mu$m, for M-type members with spectral types between M0 and M6, \\cite{sch12} derive updated excess fractions of $45^{+15}_{-13}$\\% for the $\\sim$5-8 Myr \\citep{luh04} $\\eta$ Cha cluster, and $21^{+12}_{-6}$\\% for members of the TW Hydrae Association (``TWA\" -- age $\\sim$8 Myr; \\citealt{zuck04}).  For the $\\sim$12 Myr old Beta Pictoris Moving Group (BPMG -- age $\\sim$12 Myr; \\citealt{zuck04}), no evidence was found for protoplanetary (primordial or transitional) disks.  Of 20 M-type stars in the BPMG in this spectral range, only one case of marginal 22 $\\mu$m excess was recovered, coming from a well-known debris disk bearing member AU Mic.  This implies that M-type stars with spectral types between M0 and M6 exhibiting mid-IR excess are very likely young (age $\\lesssim$ 10 Myr).  One can utilize this association of mid-IR  excess and youth as a new search method for identifying nearby, young ($<$10 Myr), late-type stars and brown dwarfs.  Low-mass stars and brown dwarfs in this age range should all show additional unambiguous indicators of youth, such as strong H$\\alpha$ emission, lithium absorption, low-gravity spectral features, etc.  Therefore, the youth of any candidate young M-type object discovered by its excess emission at mid-IR wavelengths can be evaluated with follow-up spectroscopy.  Although several young ($\\lesssim$100 Myr), nearby ($\\leq$80 pc) moving groups were identified during the past decade (\\citealt{zuck04}, \\citealt{tor08}), these include few low-mass members (spectral types later than $\\sim$M3).  This is mainly due to the fact that unambiguous identification of young M-type stars is difficult if a well-measured trigonometric parallax is missing.  Some activity indicators (e.g. H$\\alpha$ emission, X-ray emission) cannot readily discern young ($\\lesssim$100 Myr) stars from a pool of  old ($\\geq$600 Myr) field stars \\citep{shk09}.  The Li $\\lambda$6708 absorption feature, an effective age indicator for young FGK-type stars in the age range of 10-100 Myr, becomes increasingly sensitive to mass and age for M-type stars.  The age at which a star becomes hot enough to burn lithium in its core can be seen in the so-called ``lithium depletion boundary'', which can affect early to mid-M type stars at ages as young as the BPMG ($\\sim$12 Myr).  Some young M-type stars have been identified, though they are either co-moving companions of earlier spectral types or in a group of very young ($\\lesssim$10 Myr) stars that are relatively confined to a small region of the sky (e.g. TWA and the subregions of the Scorpius Centaurus Complex).  With a well-measured trigonometric parallax, one can find a young M-type field star (e.g. AP Col; \\citealt{rie11}), however, a systematic search for young M-type stars needs to wait for the next generation of parallax missions, such as {\\it Gaia}, {\\it Pan-STARRS}, etc.  Using the fact that random field M-type stars with mid-IR excess are extremely rare, one can search for $\\lesssim$10 Myr M-type stars in the solar neighborhood with WISE data.  These additional youngest M-type post T-Tauri stars can give an important clue to the mass function of young nearby stellar associations.  The TW Hydrae Association is one of the nearest (d$\\sim$30-90pc) star forming regions.  Its proximity, in combination with its young age make it an ideal area for the study of stellar, planetary, and circumstellar disk evolution.  It has also been shown to be useful as a testbed for the evaluation and implementation of new techniques to identify young stars (e.g. UV-excess; \\citealt{rod11}, \\citealt{shk11}). For the above reasons, we have chosen to examine the TW Hydrae Association in a pilot study to find young, M-type stars by their mid-IR excess emission. ",
        "conclusions": "We note that, at the age of TWA or younger, use of mid-IR excess as a search method is highly effective for TW Hydrae members with spectral types of M4 or later.  Including TWA 33 and TWA 34, there are 16 TWA members with spectral types in this range.  Fourteen of these were individually detected with WISE (TWA 3B and 5B are too close to their primary members to be resolved individually).  Nine of these fourteen members show significant mid-IR excess emission, eight of which were found with our search method.  The discovery of TWA 33 and TWA 34 demonstrates the effectiveness of using mid-IR excess as a tool to identify nearby, young, late-type stars.   These stars show many signatures of youth expected for a TWA member, such as H$\\alpha$ emission, lithium absorption, and low surface gravity spectral features.  Using two different methods, we estimate a distance of $\\lesssim$ 50 pc for both stars, consistent with their TWA membership.   We also rediscovered the young, likely LCC member 2M1337, first identified as a young star by \\cite{rod11}, and found evidence for the presence of a circumstellar disk around this star.  Considering the high fraction of late-type mid-IR excess stars from TWA and the slightly younger $\\eta$ Cha Association (45\\% for M-type members - Schneider et al. 2012), we believe that mid-IR excess can be a  a useful tool in identifying more young, nearby, late-type stars and brown dwarfs.  One can expand our search for these objects to other areas of interest, such as the sub-regions of the Scorpius-Centaurus Association, or even the entire sky.  An interesting conjecture is whether there is a higher fraction of mid-IR excess stars among the latest M spectral types (M6 or later).  For the six known TWA members with spectral type later than M6, three -- TWA 27, 30A and 32 -- were detected individually by WISE to have excess emission in channel 4.  TWA 5B, an M8 dwarf companion to TWA 5A, is too close to its primary to be resolved.  TWA 26 and TWA 29 have only upper limits in WISE channel 4, so no strong conclusions regarding mid-IR excess can be deduced from WISE for these members, though TWA 26 shows no excess at 24 $\\mu$m via Spitzer MIPS \\citep{sch12}.  In summary, all WISE detectable late M-type TWA members show mid-IR excess emission at the W4 band.  To date, no members of the $\\eta$ Cha cluster have been found with spectral types later than M6, though searches have been performed \\citep{lyo06}.  The BPMG has two known late-type members (excluding giant planet $\\beta$ Pictoris b), HR 7329B (M7.5; \\citealt{low00}) and 2MASS 0608-27 (M8.5; \\citealt{rice10}).  HR 7329B is too close to its host star to be resolved individually with WISE, and 2MASS 0608-27 has an upper limit in WISE channel 4.  So, at this point, only TWA has a reasonable number of latest M-type members to check the conjecture of even more prolonged disks among the latest M-type stars.  Future missions, such as {\\it Gaia}, which should discover more late-type members of these nearby associations, will allow us to test the hypothesis that primordial disks have longer lifetimes around later-spectral types in this spectral type range.  While we have shown that mid-IR excess emission can be a useful tool for identifying nearby, young, late-type stars and brown dwarfs, caution must be taken when considering association properties, such as disk fraction and mass function.  Any objects found with this method will surely bias the measure of disk frequency for a particular group of stars, so we do not update the excess fraction for TWA here.  If the hypothesis of longer disk lifetimes around later spectral types is true, then any objects found with this method would bias any estimate of the mass function as well.  While the question of whether or not there is a higher fraction of mid-IR excess for the latest spectral types is an interesting one, it cannot be answered with objects found with the search method described in this paper."
    },
    "1208/1208.4631_arXiv.txt": {
        "abstract": "The \\emph{Sun Watcher with Active Pixels and Image Processing} (SWAP) is an EUV solar telescope on board ESA's \\emph{Project for Onboard Autonomy~2} (PROBA2) mission launched on 2~November~2009. SWAP has a spectral  bandpass centered on 17.4~nm and provides images of the low solar corona over  a 54$\\times$54~arcmin field-of-view with 3.2~arcsec pixels and an imaging cadence of about two~minutes.  SWAP is designed to monitor all space-weather-relevant events and features in the low solar corona. Given the limited resources of the PROBA2 microsatellite,  the SWAP telescope is designed with various innovative technologies, including an off-axis optical design and a CMOS--APS detector. This article provides reference documentation for users of the SWAP image data. ",
        "introduction": "\\label{Introduction}  The \\emph{Project for On Board Autonomy} (PROBA) missions are a series of microsatellites launched by the European Space Agency (ESA) and intended to provide an in-orbit test platform for new technologies. PROBA1 was launched in 2001 and carried instruments primarily intended for Earth observation.  PROBA2 \\cite{TI_PROBA2_Platform_Santandrea} was inserted into a Sun-synchronous polar orbit at an altitude of approximately 720~km. This orbit guarantees nearly uninterrupted solar viewing through most of the year. The nominal two-year mission started with the 2~November~2009 launch but was subsequently extended until the end 2012.  The primary mission goal of PROBA2 is to perform an in-flight demonstration of a series of new spacecraft technologies. The secondary mission goal is the exploitation of the payload of scientific instruments consisting of two Sun-sensing instruments, the \\emph{Sun Watcher with Active Pixel Sensor and Image Processing} (SWAP), the main subject of this article, and the \\emph{Lyman-Alpha Radiometer} (LYRA: \\opencite{Hochedez06}, \\opencite{TI_PROBA2_LYRAinstrumentPaper_Dominique}).  Also on board are two instruments intended to measure {\\it in-situ} properties of the magnetosphere: the \\emph{Thermal Plasma Measurement Unit} (TPMU) and the \\emph{Dual-Segmented Langmuir Probe} (DSLP).  \\begin{sloppypar} As a microsatellite, the PROBA2 platform imposes severe limitations on the payload resources such as volume, mass, and power. Thus SWAP was designed as a miniaturized coronal imager, capable of meeting these limitations, but also capable of extending the successful ``CME watch'' program that began with data from the \\emph{Extreme-Ultraviolet Imaging Telescope} (EIT) on board the joint ESA--NASA \\emph{Solar and Heliospheric Observatory} (SOHO) mission \\cite{Delaboud95}. By improving the imaging cadence by an order of magnitude over that of EIT, SWAP can track the evolution of all events in the low solar corona that are relevant for space weather, including flares, CMEs, EUV waves, and EUV dimmings. In addition, SWAP continuously images solar features such as coronal holes and active regions, the locations of which are essential data for space weather forecasters. As a result, while SWAP provides useful data for both scientific analysis and space weather forecasting in the present, the lessons learned during the development of SWAP are being incorporated into a future generation of space weather monitors as well. One example is the \\href{http://sidc.be/esio/}{\\emph{EUVI Solar Imager for Operations} (ESIO)}, a potential building block of the ESA Space Situational Awareness Program. \\end{sloppypar}  Compared to more resource-rich contemporary EUV imagers such as the \\emph{Extreme Ultraviolet Imagers} (EUVI), part of the \\emph{Sun Earth Connection Coronal and Heliospheric Investigation} (SECCHI) packages on the twin \\emph{Solar Terrestrial Relations Observatory} (STEREO) spacecraft \\cite{2008SSRv..136...67H}, and the \\emph{Atmospheric Imaging Assembly} (AIA) instrument on board the \\emph{Solar Dynamics Observatory} (SDO; \\opencite{2012SoPh..275...17L}), SWAP offers only modest temporal and spatial resolution. However, SWAP's design and operational strategy result in some unique capabilities compared to other EUV solar imagers.  In particular, SWAP images are useful because they provide the largest EUV field-of-view available from Earth orbit (54~arcmin) and SWAP's complementary metal-oxide-semiconductor active-pixel system (CMOS--APS) detector does not bloom significantly during bright flares, when nearly all other detectors do so. In addition, the agility feature of PROBA2 permits SWAP to off-point by up to one~degree, allowing imaging of extended coronal features and tracking of CMEs as they propagate away from the Sun.  In this article we provide a reference description of the SWAP instrument, the results of several pre-flight instrumental tests, and an overview of spacecraft operations that are likely to be particularly relevant to data users in the solar physics and space weather communities. The second part of this article, \\inlinecite{TI_PROBA2_SWAPcalibrationPaper_Halain}, discusses in-flight calibration and performance of the SWAP instrument. \\inlinecite{Berghmans06} discuss the early development of the SWAP instrument, while many of the technological improvements associated with the development of SWAP are reviewed by \\inlinecite{Defise07}.  The material presented here is a synthesis of technical specifications that appear in the data sheets for the several instrument and spacecraft subsystems, the instrument design itself, and from data collected during two calibration campaigns at the \\href{http://www.ptb.de/mls/aufgaben/bessylab.html}{PTB/BESSY} facilities in Berlin, Germany, in February 2007 and July 2008. These calibration campaigns included tests of individual parts of the SWAP instrument as well as the complete SWAP optical path and detector. (During the July 2008 campaign we conducted these tests using the flight models of the instrument proximity electronics and instrument interface unit.)  In Section~\\ref{Overview} we give an overview of the instrument and introduce the main subsystems of the telescope. In the following sections we focus on the main functional aspects of the SWAP telescope that are relevant for SWAP data users: spectral selection (Section~\\ref{Spectral}), imaging (Section~\\ref{Imaging}), signal recording (Section~\\ref{SignalRecording}), and on board data processing (Section~\\ref{OnboardDataProcessing}).  After that we discuss the limitations and opportunities provided by SWAP's operational strategy (Section~\\ref{Operations}) as well as the SWAP data products and analysis software that are available to users (Section~\\ref{DataProducts}) through the PROBA2 Science Center. Finally, we make some concluding remarks in Section~\\ref{Conclusions}. For quick access to reference material, a SWAP data sheet is provided in an appendix.  Additional articles in this same \\emph{Solar Physics Topical Issue} concentrate on the PROBA2 platform \\cite{TI_PROBA2_Platform_Santandrea}, on the LYRA ``sister'' instrument \\cite{TI_PROBA2_LYRAinstrumentPaper_Dominique}, and on the management of SWAP and LYRA operations from the PROBA2 Science Center \\cite{TI_PROBA2_P2SC_Zender}. ",
        "conclusions": "\\label{Conclusions}   The experience gained from building and operating the SWAP instrument is of particular relevance for developers of the next generation of solar imagers.  SWAP's small size and minimal power requirements have proven that successful science can be carried out even using instruments that operate on resource-poor platforms. SWAP's on board processing capabilities, which both reduce the telemetry required to download data from the spacecraft and optimize the use of limited on board storage by identifying images of maximal interest, provide a foundation for operational planning for the three-bandpass \\emph{Extreme-Ultraviolet Imager} (EUI) on the proposed ESA \\emph{Solar Orbiter} mission. Given the highly limited telemetry bandwidth available for the mission, on board processing and storage of images is of the utmost importance.  Additionally, EUI, which combines two co-aligned \\emph{High Resolution Imagers} (HRI) and one \\emph{Full Sun Imager} (FSI) will make use of SWAP's optical design principles as well: one of the HRIs is a two-mirror off-axis telescope with the same bandpass as that of SWAP.  SWAP has also demonstrated that CMOS--APS detectors, which offer advantages of traditional CCD-based detectors, are fully ready for use in scientific remote sensing applications. It is likely that several upcoming solar-imaging missions will make use of similar detectors, and the lessons learned from the SWAP detector can help to inform design and planning for these future instruments.  Finally, with its large field-of-view\\,---\\,the largest EUV field-of-view currently available from the perspective of the Earth\\,---\\,and relatively high cadence, SWAP provides images that are simultaneously useful for both science applications and for space-weather monitoring and forecasting. Lessons learned from SWAP have also informed the design of future space-weather monitors, which, given the low overhead required, could easily be integrated into any space-based platform with a small amount of additional space and power.     \\appendix   For quick user reference, Table~\\ref{instchar_tab1} provides a summary of the main properties of SWAP that are relevant to general users.  \\begin{table} \\caption{Instrument Data Sheet} \\label{instchar_tab1} \\begin{tabular}{l c c c} \\hline\\hline Property & Parameter & Value & Units \\\\ \\hline Wavelength of peak response \t\t\t& $\\lambda_{\\mathrm{max}}$ & 17.4 & nm \\\\ Quantum efficiency of detector \t\t& $QE$ & 0.45 & n/a \\\\ e$^{-}$ per DN conversion factor \t& $C_{e^{-}/DN}$ & 31 & n/a \\\\ Geometric aperture area \t\t\t\t& $A_{\\mathrm{aper}}$ & 8.55 & cm$^{2}$ \\\\ Geometric pixel size \t\t\t\t& $l_{\\mathrm{pix}}$ & 0.018 & mm \\\\ Telescope focal length \t\t\t\t& $f$ & 1173 & mm \\\\ Angular pixel size \t\t\t\t\t& $l_{\\mathrm{pix}}$/$f$ & 3.17 & arcsec \\\\ Angular coverage per pixel \t\t\t& ($l_{\\mathrm{pix}}$/$f$)$^{2}$ & $2.35 \\times 10^{-10}$ & str \\\\ \\hline \\end{tabular} \\end{table}           \\begin{acks}  SWAP is a project of the Centre Spatial de Li\\`ege and the Royal Observatory of Belgium funded by the Belgian Federal Science Policy Office (BELSPO). The work at these institutes and at the K.U.Leuven is supported by PRODEX grant C90193 (SWAP - Preparation to Exploitation), managed by the European Space Agency in collaboration with the Belgian Federal Science Policy Office (BELSPO). We also acknowledge the financial support of the Solar Terrestrial Center of Excellence during the design, testing, and operations of SWAP. We gratefully acknowledge many helpful discussions with  J.-F. Hochedez during the design and testing of SWAP. We acknowledge financial support from the Max-Planck-Institut f\\\"ur Sonnensystemforschung, Germany, for making the SWAP calibration campaign possible through collaboration with Physikalisch-Technische Bundesanstalt (PTB). We also thank the PTB team of Frank Scholze for their dedication and support during the SWAP calibration campaigns. Part of this work was supported by the German \\emph{Deut\\-sche For\\-schungs\\-ge\\-mein\\-schaft, DFG\\/} project number Ts~17/2--1. \\end{acks}"
    },
    "1208/1208.6544_arXiv.txt": {
        "abstract": "We use radiative transfer to study the growth of ionized regions around the brightest, $z=8$ quasars in a large cosmological hydrodynamic simulation that includes black hole growth and feedback (the MassiveBlack simulation). We find that in the presence of the quasars the comoving HII bubble radii reach $\\unit[10]{Mpc/h}$ after 20 Myear while with the stellar component alone the HII bubbles are smaller by at least an order of magnitude. Our calculations show that several features are not captured within an analytic growth model of Stromgren spheres.  The X-ray photons from hard quasar spectra drive a smooth transition from fully neutral to partially neutral in the ionization front. However the transition from partially neutral to fully ionized is significantly more complex. We measure the distance to the edge of bubbles as a function of angle and use the standard deviation of these distances as a diagnostic of the anisotropy of ionized regions. We find that the overlapping of nearby ionized regions from clustered halos not only increases the anisotropy, but also is the main mechanism which allows the outer radius to grow. We therefore predict that quasar ionized bubbles at this early stage in the reionization process should be both significantly larger and more irregularly shaped than bubbles around star forming galaxies. Before the star formation rate increases and the Universe fully reionizes, quasar bubbles will form the most striking and recognizable features in 21cm maps. ",
        "introduction": "The current consensus is that the contribution to the global ionizing budget from quasars during the Epoch of Reionization (EoR) is small compared to that from early stars and galaxies \\citep[see, eg, ][]{2009JCAP...03..022L,1996ApJS..102..191G,2011ASL.....4..228T}. The EoR began as the population III stars and primordial galaxies ionized their most immediate vicinity, as studied by \\citet{2008MNRAS.384.1080T, 2008ApJ...684...18C, 2011MNRAS.417.2264V}. On the other hand, the characteristic proper radius of ionizing bubbles at the end of EoR is constrained to be on the order of $\\unit[10]{Mpc/h}$ \\citep{2007MNRAS.376L..34M, 2007MNRAS.380L..30A, 2004Natur.432..194W}, and the quasar contribution is limited to be less than 14\\% \\citep{2007MNRAS.374..627S} of the total.  Even though the contribution to global reionization by quasars is constrained in this way, bright quasars may still leave a signature on the growth of individual ionized regions.  Several authors have investigated this signature in mock 21cm emission spectra taken from simulations of isolated quasars \\citep{2012MNRAS.424..762D, 2011arXiv1111.6354M, 2010ApJ...723L..17A, 2008MNRAS.391.1900D, 2008MNRAS.386.1683G} at redshift $z \\sim 8$.  In observations, an object recently reported by \\citet{2011Natur.474..616M}, ULAS J1120+0641, has a luminosity of $6.3\\times10^{13}L_{\\odot}$ at $z \\sim 7$ and a proper near-zone radius of less than $\\unit[2]{Mpc/h}$.  The near-zone radius is consistent with the possibility of bright quasar driven growth in a near neutral intergalactic medium background \\citep{2011MNRAS.416L..70B}.  In this paper we study the growth of the ionization front of bright quasars in an almost neutral cosmological context.  The quasars and their surrounding medium are selected from a large hydrodynamic simulation \\citep[the MassiveBlack simulation, introduced in][]{2012ApJ...745L..29D}, and then post-processed with a radiative transfer code. This allows us to simulate 8 rare quasars using reasonable computing resources.  Our focus is on the evolution and properties of the largest individual ionized bubbles, the sources that produce them, and the relationship between the two.  Because the simulation forms quasars and star forming galaxies ab initio, we are able to make use of the luminosity and positions of the radiation sources that the simulation produces, rather than setting them in by hand. However we do not deal with the full reionization of the volume of the simulation, which would require following the evolution of the entire density and ionization field from high redshifts down to at least $z=6$.  Instead we restrict ourselves to the growth of ionized regions around an early period in this process (at $z=8$), where the photon path lengths are still smaller than the computational sub-volumes we analyze. We leave the study of the full reionization of the volume to future work. ",
        "conclusions": "We have presented results from radiative transfer simulations in the vicinity of high redshift quasars in the MassiveBlack simulation.  We find that the rare brightest quasars drive a much more significant growth of ionized regions than in the purely stellar driven case. The ionized regions associated with active quasars are characterized by (i) a smooth ionized fraction transition from the middle to the outer front, and (ii) an increased anisotropy in the front when it starts to overlap the nearby ionized regions.  The nature of such growth is significantly more complex than a simple analytic growth of a single center bubble with clumping correction.  The largest HII bubble obtained in this simulation has a comoving radius of $\\unit[10]{Mpc/h}$, which is smaller than the general expectation that can fulfill the reionization of the universe \\citep{2011ASL.....4..228T, 2004Natur.432..194W, 2011arXiv1111.6354M}. The quasar near zones we have presented in this paper are the primordial ancestors of the later much larger Stromgren spheres which will form near the end of the EoR. They are however relatively isolated regions that could be interesting objects for study in future 21cm surveys. After the $z=8$ epoch we have modeled in this paper, we expect that the global star formation in the MassiveBlack simulation increases significantly. This will eventually lead to the global reionization of the universe. We plan to study this process and role of quasars in future work.  \\textbf{Acknowledgments:} We acknowledge support from Moore foundation which enabled the radiative transfer simulations to be run at the McWilliams Center of Cosmology at Carnegie Mellon University. The MassiveBlack simulation was run on the Cray XT5 supercomputer Kraken at the National Institute for Computational Sciences. This research has been funded by the National Science Foundation (NSF) PetaApps program, OCI-0749212 and by NSF AST-1009781.  We also acknowledge the support of a Leverhulme Trust visiting professorship. The visualizations were produced using Healpix \\citep{2005ApJ...622..759G}, Healpy \\footnote{http://github.com/healpy/healpy}, and Gaepsi \\footnote{https://github.com/rainwoodman/gaepsi} \\citep{2011ApJS..197...18F}."
    },
    "1208/1208.5535.txt": {
        "abstract": "We present a full analysis of the Probing Evolution And Reionization Spectroscopically (PEARS) slitess grism spectroscopic data obtained with the Advanced Camera for Surveys onboard {\\it HST}. PEARS covers fields within both the Great Observatories Origins Deep Survey (GOODS) North and South fields, making it ideal as a random survey of galaxies, as well as the availability of a wide variety of ancillary observations complemented by the spectroscopic results.  Using the PEARS data we are able to identify star-forming galaxies within the redshift volume 0 $<$ z $<$ 1.5.  Star-forming regions in the PEARS survey are pinpointed independently of the host galaxy.  This method allows us to detect the presence of multiple emission line regions (ELRs) within a single galaxy. We identified a total of 1162  \\Ha, \\OIII\\ and/or \\OII\\ emission-lines in the PEARS sample of  906 galaxies to a limiting flux of $\\sim$ ${\\rm 10^{-18}\\ erg/s/cm^2}$.   The ELRs have also been compared to the properties of the host galaxy, including morphology, luminosity, and mass.  From this analysis, we find three key results: 1)  The computed line luminosities show evidence of a flattening in the luminosity function with increasing redshift; 2) The star-forming systems show evidence of complex morphologies, with star formation occurring predominantly within one effective (half-light) radius. However, the morphologies show no correlation with host stellar mass;  and 3) The number density of star-forming galaxies with {\\it M}$_{*}$ $\\geq$ 10$^{9}$ {\\it M}$_{\\odot}$ decreases by an order of magnitude at z $\\leq$ 0.5 relative to the number at 0.5 $<$ z $<$ 0.9 supporting the argument of galaxy downsizing. ",
        "introduction": "%NOTE:  There needs to be some additional information about the Madau Plot and %morphological evolution here. Maybe 1 paragraph each since these relate directly %to the main results. \\indent Emission line galaxies (ELGs) are systems selected by the presence of strong line emissions (e.g., \\Lya\\, \\OII, \\OIII, \\Hb, and \\Ha), usually detected using narrow band  or grism surveys.  The strong emission lines in these galaxies trace recent star formation activity, in contrast to the star formation history and properties of the global stellar populations that can be discerned using broad band observations.  A significant amount of the light originating from ELGs is contained in regions producing strong emission lines, which in turns makes these objects easily identifiable. The emission lines of ELGs also provide a convenient way to determine the redshifts of these objects. Since ELGs are selected on the basis of strong emission lines, rather than continuum emission, selecting ELGs allows one to potentially probe down to a  lower luminosity --- and thus lower mass -- galaxies compared to broad band surveys, which tend to be limited by the luminosity of the galaxies themselves, rather than the strength of their emission  lines.   Assuming that ELGs are spatially distributed in a fashion similar to other galaxies, they provide a powerful tool for tracing the star formation history of the Universe \\citep[e.g.,][]{Salzer1988,Popescu1997}.\\\\  \\indent The epoch $0 < z < 1.5$  discussed in this paper is important, because star formation activity in galaxies has been observed to increase significantly  as redshift increases \\citep[e.g.][]{Madau1998,hopkins2004}. While at higher redshifts ($z>2$), there is still some controversy as to whether the star-formation rate density (SFRD) relation flattens or decreases, the initial increase in star-formation  implies that, at low z, some mechanism(s) must have occurred, which quenched star formation. If this was not the case, massive ellipticals today would still be observed to be strongly forming many stars. There is also evidence that suggest that the interstellar medium, star formation rates and gas fractions differ between local and distant galaxies. Studying galaxy evolution at these redshifts therefore does not only require the ability to measure the star formation history of these objects, but also depends on our ability to properly sample galaxies over a wide range of masses to alleviate as many biases as possible. ELGs are ideal for such work.  As noted above, these objects are easily detected in surveys and they are efficient for probing to lower stellar masses in terms of telescope time required.  The wavelength range of the ACS grism used for PEARS makes it possible to identify the strong rest-frame emission lines that are well known to be a sign of vigorous star formation (e.g. \\Ha, \\OIII\\ and \\OII) out to z $\\simeq$ 1.5.  In this paper,  examining  \\Ha, \\OIII\\ and \\OII\\  emitters separately allows us to look at properties of star-forming galaxies in increasing redshift ranges.  When plotted separately, these three emission lines represent proxies for the redshift bins of $0<z<0.5$, $0.1<z<0.9$, and $0.5<z<1.5$, respectively.\\\\  \\indent Identifying ELGs has traditionally been done using narrow-band photometric filters.  This technique has also been successfully applied to very high redshifts to detect \\Lya\\ emitters \\citep{Rhoads2001}. However, while narrow-band surveys can efficiently cover large fields-of-view to relatively faint magnitudes, they are typically limited to very small and discrete redshift ranges.  This can be partially alleviated using multiple narrow band filters \\citep[e.g. See Subaru Deep Field,][]{Ly2007}, but continuous redshift coverage remains intrinsically limited in narrow-band surveys. The Probing Evolution And Reionization Spectroscopically (PEARS) slitless grism spectroscopic survey provides an unprecedented opportunity to study ELGs in a way that cannot be achieved from any ground-based observations.  PEARS allows us to bypass the difficulties inherent in narrow-band filter surveys (as noted above) and the limitations imposed by varying sky-brightness and atmospheric emission lines, which can limit ground-based grism surveys, and identify ELGs based {\\it solely} on the direct detection of emission lines in dispersed slitless spectra. As previously shown \\citep{pirzkal2006,straughn2008,straughn2009}, this approach allows us to detect emission lines in very faint host galaxies, particularly sub-{\\it m}$^{*}$ galaxies, over a very large and continuous redshift range. Since our survey is mainly limited by the line fluxes themselves, faint emission lines can be identified in galaxies that are only weakly detected (i.e. high EW emission lines) while brighter host galaxies tend to increase the local background flux, diluting fainter emission lines in the brightest hosts. The line flux limit of this survey is discussed in more details in Section \\ref{completeness}. Particular to the PEARS survey, the use of multiple position angles on the sky allows us to identify emission lines using independent observations, as well as to pinpoint the exact physical location of the ELRs within each ELG. Crucially, and in addition to this, the PEARS survey was designed to overlap with both the GOODS-N and GOODS-S fields, so that there exists a substantial amount of ancillary data available, including very deep, high-resolution broad-band imaging ranging from the UV to the infrared bands.  As we noted above, the redshift range (0 $<$ z $<$ 1.5) probed by PEARS is a critical transition epoch, both in terms of star-formation histories and morphological evolution. On one hand, the PEARS grism slitless observations make it possible to efficiently identify emission lines, identify the corresponding ELRs and host ELGs.  On the other hand, the GOODS ancillary data allow us to examine the morphology and physical characteristics of the ELGs. This powerful combination of data gives us an opportunity to examine the evolution of ELGs over a long period of cosmic time and over a much wider mass range than has been previously probed.\\\\  \\indent This paper is organized as follow: Section \\ref{observations} briefly summarizes the PEARS observations (HST Proposal 10530, P.I. Malhotra). Section \\ref{analysis} describes the data reduction and analysis of the sample, including detection, extraction and identification of emission lines, as well as completeness tests. Section \\ref{ELR} presents the PEARS  \\OII, \\OIII\\ and \\Ha\\ line luminosity functions and their redshift evolution. Finally, Section \\ref{hosts} compares the properties of the PEARS host galaxies, such as morphology and luminosity, with the star formation properties discerned from the PEARS emission lines. All calculations in this paper assume $H_0={\\rm 70\\ km\\ s^{-1}\\ Mpc^{-1}}$ and $\\Omega_M  = 0.3$, $\\Omega_\\Lambda = 0.7$ \\citep{Komatsu2011,Hinshaw2012}. All magnitudes are given in the AB system. ",
        "conclusions": "\\indent We have presented a sample of ELGs selected independently by their emission-lines without {\\it a priori} knowledge of their host galaxies properties.  The methodology used (PEARS-2D) is based on {\\it direct} detection of emissions-line from {\\it HST} slitless grism spectroscopy, with the added bonus of being able to detect multiple ELRs within a single galaxy.  This  has yielded  a sample, which is effectively random and blind to other parameters.  Using the wealth of ancillary data, the properties of the underlying host galaxies were compared with the SFR histories derived from the ELRs.  The key results are summarized below:\\\\ \\noindent 1)  There is evidence for evolution in the luminosity functions of  the \\Ha,\\OIII\\ and \\OII\\ emission lines. The luminosity function slopes flatten as a function of redshift.\\\\ \\noindent 2)The morphology of the host galaxies clearly indicates that these objects are clumpy, although we detect no correlation between their morphology and our stellar mass estimates or star-formation intensity (sSFR).\\\\ \\noindent 3)The mass-density function of \\OIII\\ emitting galaxies at ${\\rm 0<z<0.9}$ strongly decreases with redshift. The number density of objects with stellar masses greater that ${\\rm \\sim 10^{10}\\ M_{\\sun}}$ undergoing strong star-formation decreases at lower redshifts. This supports the idea of galaxy downsizing \\citep[][e.g.]{cowie1996}.\\\\ \\noindent  The results presented here also demonstrate the clear advantage of space-based grism spectroscopy using multiple position angles.  Such observations are able to probe deeper than similar ground-based studies.  The PEARS-2D method also provides a method for detecting multiple ELRs, and allows spatial information about SF to be derived for galaxies.  Future work will include using the WFC3 near-IR grism mode, with observed wavelength coverage of 0.8-1.6$\\micron$.  This will allow us to probe to significantly higher redshifts, and determine whether the trends reported here continue to earlier epochs. After 2018, this work can also be done at much higher sensitivities with the JWST FGS grism at $1-2 \\mu m$, and the JWST NIRcam prisms at $2.5-5 \\mu m$."
    },
    "1208/1208.3682_arXiv.txt": {
        "abstract": "We address the effects of bar-driven secular evolution in discs by comparing their properties in a sample of nearly 700 unbarred and barred (42 $\\pm$ 3 per cent of the population) massive  disc galaxies (\\mstar\\ $\\ge 10^{10}$ \\msun). We make use of accurate structural parameters derived from $i$-band bulge/disc/bar decompositions to show that, as a population, barred discs tend to have fainter central surface brightness ($\\Delta$\\muo\\ $\\approx 0.25$ mag), and disc scale lengths that are $\\approx$\\,15 per cent larger than those of unbarred galaxies of the same stellar mass. The corresponding distributions of \\muo\\ and $h$ are statistically inconsistent at the $5.2\\,\\sigma$ and $3.8\\,\\sigma$ levels, respectively. Bars rarely occur in high-surface brightness discs, with less than 5 per cent of the barred population having \\muo\\ $< 19.5$ \\sb\\ -- compared to 20 per cent for unbarred galaxies. They tend to reside in moderately blue discs, with a bar fraction that peaks at $(g-i)_{disc} \\approx 0.95$ mag and mildly declines for both bluer and redder colours. These results demonstrate noticeable structural differences between the discs of barred and unbarred galaxies, which we argue are the result of bar-driven evolution -- in qualitative agreement with longstanding theoretical expectations. ",
        "introduction": "\\label{sect:intro}  Secular processes are expected to have played a significant role in establishing the current properties of massive disc galaxies. These are dynamically evolved systems, as evidenced by the presence of a population of discs at $z\\sim1$ with similar scale lengths \\citep{Lilly1998,Simard1999,Barden2005} and bar fraction \\citep{Jogee2004,Sheth2008} to what is found in the Local Universe. This suggests that the quiescent phase of \\emph{massive} discs evolution started at least 8 Gyr ago, thus leaving ample time for secular mechanisms to operate (but see \\citealt{Hammer2005}).  Bars perhaps provide the most clear-cut evidence of the impact of secular evolution on disc galaxies. Analytical and numerical calculations show that they drive a significant amount of mass and angular momentum redistribution in the disc \\citep{Hohl1971,Sellwood1993,Athanassoula2003,Martinez-Valpuesta2006}, funneling material towards the galaxy inner regions that can result in enhanced gas and stellar densities -- thus possibly fueling nuclear starbursts and AGNs \\citep{Shlosman1989} and leading to the formation of central bulge-like components \\citep{Courteau1996,Kormendy2004,Debattista2006}. Furthermore, the angular momentum exchange between the inner and outer disc can lead to increased disc scale lengths, and the development of mass profile breaks at large radii \\citep{Valenzuela2003,Debattista2006}.  It is precisely the amount of angular momentum exchanged within the galaxy what determines the bar strength, and this ultimately depends on the mass and velocity distributions of the material in the disc and spheroidal (bulge plus halo) components \\citep{Athanassoula2003}. Once formed, these non-axisymmetric perturbations appear to be long-lived, surviving buckling instabilities and the growth of a significant central mass concentration ($>$\\,10\\%, and perhaps up to 20\\% of the original disc mass; \\citealt{Shen2004,Athanassoula2005,Debattista2006}). On the other hand, \\citet{Bournaud2005} suggest that a substantial gas component in the disc ($\\sim$\\,7\\% of the total visible mass) might lead to the weakening, or even destruction of the bar due to the angular momentum exchange resulting from gas inflow. However, this effect is significantly reduced when a responsive dark matter halo model is considered \\citep{Berentzen2007}.  The importance of bars in galaxy evolution studies does not only stem from the fact that they can significantly impact the evolution of a given galaxy, but also because such structures are rather common in disc galaxies. The most recent optical studies indicate that approximately half of all massive disc galaxies contain bars \\citep{Barazza2008,Aguerri2009}. The fraction of \\emph{strong} bars increases substantially in dust-penetrating near-infrared wavelengths ($\\approx$\\,60 per cent, \\citealt{Eskridge2000}), while the \\emph{total} bar fraction does so by up to 35 per cent, either when bar detection is carried out through visual classification \\citep{Buta2010} or via quantitative methods \\citep{Marinova2007,Menendez-Delmestre2007,Weinzirl2009}. It is now well established, though, that the bar fraction in the Local Universe is a strong function of the galaxy stellar mass (\\citealt{Mendez-Abreu2010,Mendez-Abreu2012,Nair2010,Cameron2010}), and thus consistent results can only be obtained when comparing samples well matched in \\mstar. Even though wavelength coverage, detection technique, and --possibly to a larger extent-- sample selection, all play a role in the detection of bars in disc galaxies, most studies point to fractions $\\sim$\\,50 per cent in $L \\gtrsim L^{*}$ galaxies.  The picture that emerges from theoretical work is corroborated, at least qualitatively, by a number of observational results. The bar-driven redistribution of angular momentum affects the interstellar medium in galaxies, resulting in flatter chemical abundance gradients in barred galaxies \\citep{Martin1994,Zaritsky1994}, as well as higher central concentrations of molecular gas \\citep{Sakamoto1999,Sheth2005}. More recently, evidence has been found that the gas brought to the centre by bars is efficiently transformed into stars. \\citet{Ellison2011} present indication that the current star formation rate at the centre is higher in massive barred galaxies. \\citet{Coelho2011} show that the distribution of mean stellar ages in bulges of massive barred galaxies shows a peak at low ages that is absent for their unbarred counterparts \\citep[see also][]{Perez2011}.  Yet investigation of the effects of bar-driven secular evolution on the structural properties of discs is currently lacking in the literature. In this Letter we fill this gap by comparing the disc properties in sample of nearly 700 barred and unbarred galaxies. ",
        "conclusions": "\\label{sect:discussion}  We have shown that barred and unbarred \\mstar\\ $\\ge 10^{10}$ \\msun\\ galaxies are characterised by distinct distributions of disc structural parameters. As a population, barred discs tend to have fainter central surface brightness ($\\Delta$\\muo\\ $\\approx 0.25$ mag in the $i$-band) and disc scale lengths that are larger by $\\approx$\\,15 per cent than those of unbarred galaxies. We argue here that these differences are the result of bar-driven disc evolution. While the alternative scenario --i.e., that the distinction arises due to initial disc conditions that favoured bar formation-- is difficult to rule out solely based on observations, theoretical work strongly suggest that is not the case. First, all numerical simulations show that the onset of the bar instability results in structural evolution of the disc (to varying degrees). Given that we are studying barred galaxies, it is reasonable to expect that the disc has actually evolved, i.e., that its current properties are not exactly the same ones that led to the formation of the bar. Second, and most important, the alternative scenario would imply that bar formation is favoured in discs having lower surface brightness and larger scale lengths. This is in disagreement with the numerical results by \\citet{Mayer2004}, where it is shown that low surface brightness discs are generally rather stable against bar formation due to a combination of low self-gravity and a high halo/disc mass ratio.  The observed differences in disc scale lengths are close to, but slightly lower than,  those predicted by the numerical simulations of \\citet{Valenzuela2003}, where discs were found to increase their $h$ by factors 1.2$-$1.5 due to the transfer of angular momentum that accompanies the formation of a bar. Nevertheless, the moderate increment of scale lengths  revealed by the data has to be thought of as a lower limit to the actual effects of bar-driven secular evolution. The reason is twofold. First, it is important to recall that bars are not the only drivers of disc secular evolution.  In fact, any type of non-axisymmetric perturbation --including spiral arms, oval distortions and triaxial dark matter haloes-- can also modify the properties of discs (e.g., \\citealt{Kormendy2004,Sellwood2010}), but they generally operate on longer timescales. It is therefore possible that secular evolution has also occurred, to some degree, in our sample of unbarred galaxies. Second, our disc structural parameters come from fitting one single exponential to the surface brightness profile, with no distinction between different disc Freeman types \\citep{Freeman1970}. This can slightly bias our recovered parameters, resulting in marginally brighter \\muo\\ and smaller $h$ in the case of Type\\,II profiles with outer truncations -- which occur far more often than single exponential, Type\\,I profiles \\citep[e.g.,][]{Pohlen2006}. This is indeed supported by the recent analysis of disc properties in S$^{4}$G galaxies \\citep{Sheth2010}. Mu\\~noz-Mateos et al. (in preparation)  find that, when allowing for outer profile breaks, disc scale lengths in barred galaxies are, on average, a factor $\\sim$\\,1.8 larger than those without bars -- but, as in our case, the scatter at fixed 3.6\\,$\\mu$m magnitude is large.  In any case, our fits provide the optimal average exponential profile between small and large radii, and thus allow for a direct comparison with the numerical simulations by \\citet{Debattista2006}. Not surprisingly, they find that the amount of evolution in their simulated discs depends critically on the \\emph{initial} disc kinematics. The density profiles of highly unstable discs (low Toomre's $Q$) evolve dramatically compared to more stable initial configurations, resulting in final scale length differences of factors $\\gtrsim$\\,2 even for models with nearly identical initial angular momentum. The general situation is of course more complex than this, and the specific density profile evolution in their simulations is determined by the phase-space distributions of the stellar disc and  the dark matter halo -- resulting in scale lengths changing by factors 1.0$-$2.4. As \\citet{Debattista2004,Debattista2006} point out, this has the important consequence that direct estimates of dark matter halo spin parameters from measured disc scale lengths can be rather uncertain.  Our results support their claims, but suggest that the amount of scale length evolution due to bar formation is only moderate. A simple back-of-the-envelope argument appears to be consistent with this idea. If secular mechanisms were to increase disc scale lengths by a considerable amount --say, factors larger than two--, we would then expect to see clear signs of  evolution in the mass-size relation of discs between redshifts $0 < z < 1$. Yet observations support mild to no evolution at all \\citep{Lilly1998,Simard1999,Barden2005} and this, in turn, is consistent with real discs being reasonably stable ($Q \\gtrsim 1.5$; e.g., \\citealt{Kregel2005}). If this is the case, and considering all the assumptions involved, it is most likely that the uncertainties derived from mapping halo spin to disc scale lengths are not dominated by secular evolution effects -- but they certainly are a contributing factor.  We note here that we have also looked for any existing correlation between the bar strength, as measured by its ellipticity, and the central surface brightness and scale lengths of the discs. No trend whatsoever is found between $\\epsilon_{bar}$ and $\\mu_{0}$. While the mean bar ellipticity changes from $\\langle \\epsilon_{bar} \\rangle \\approx 0.6$ at $h\\sim2$ kpc to $\\langle \\epsilon_{bar} \\rangle \\approx 0.7$ at $h\\sim5$ kpc, the 0.1 scatter at fixed $h$ renders any potential trend statistically insignificant (a Spearman coefficient of only 0.3). The lack of significant correlations between these quantities suggests that the impact of bars on galaxy discs is rather complex, and does not solely depend on bar strength. This is in line with theoretical work indicating that a number of factors play an important role in the formation and evolution of bars.  Indeed, it is not yet clear if there exists one single condition determining whether a galaxy will be bar-stable or not. \\citet{Athanassoula2008} shows that simple criteria that do not fully capture the complexity underlying bar formation [e.g., the \\citet{Efstathiou1982} criterion] generically fail to  correctly predict  disc stability. Instead, numerous factors come into scene in order to stabilise a disc against non-axisymmetric perturbations. Thus, galaxies with very weak or no bars must have either a kinematically hot disc and/or a significant central mass concentration and/or a very low relative disc mass and/or be embedded in a quite unresponsive dark matter halo \\citep{Athanassoula2003,Sellwood2010}. From the disc component point of view, a high stellar velocity dispersion can provide significant stabilisation and prevent the formation of a bar for over a Hubble time \\citep{Athanassoula1986}. Moreover, \\citet{Sellwood2001} show that a steeply rising inner rotation curve is sufficient to bar-stabilise a disc, regardless of the dark matter content. In this context, the realisation in Fig.\\,\\ref{fig:scalings}$c$ that compact, high-surface brightness discs do \\emph{not} host bars --or if they do, they are rather small-- is intriguing, and it is tempting to think of these systems as having a high rate of shear at the centre capable of (almost) fully stabilising the disc. Kinematical studies of these galaxies are highly desirable in order to test this hypothesis and clarify their lack of bars.  Finally, basic stability criteria and numerical simulations suggest that bars are more easily triggered in gas-rich (i.e., cold), star-forming discs. This provides a natural explanation for the preferential occurrence of bars in moderately blue discs, but the exact shape of the bar fraction distribution deserves further investigation. Even though such an analysis is beyond the scope of this Letter, one can naively think that gas consumption through star formation leads to redder and kinematically hotter (more bar-stable) discs. In this context, the declining fraction towards bluer colours has probably to be understood in terms of a selection effect, such that these galaxies contain smaller bars that remain undetected in our images (cf. \\citealt{Nair2010}).  We have shown that  the discs of barred and unbarred galaxies are characterised by distinct structural properties, which we argue are the result of bar-driven evolution. Detailed structural decomposition of galaxies provides one of the most powerful diagnostics for galaxy evolution studies, allowing for direct, quantitative comparisons with theory and  simulations."
    },
    "1208/1208.0362_arXiv.txt": {
        "abstract": "After eight months of continuous observations, the Wide-field Infrared Survey Explorer (WISE) mapped the entire sky at 3.4 $\\mu$m, 4.6 $\\mu$m, 12 $\\mu$m and 22 $\\mu$m. We have begun a dedicated WISE High Resolution Galaxy Atlas (WHRGA) project to fully characterize large, nearby galaxies and produce a legacy image atlas and source catalogue. Here we summarize the deconvolution technique used to significantly improve the spatial resolution of WISE imaging, specifically designed to study the internal anatomy of nearby galaxies. As a case study, we present results for the galaxy NGC\\,1566, comparing the WISE super-resolution image processing to that of \\Spitzer, GALEX and ground-based imaging.  The is the first paper in a two part series;  results for a much larger sample of nearby galaxies is presented in the second paper. ",
        "introduction": "For nearly three decades now, starting with IRAS in the early 1980's and continuing today with the \\Spitzer and the AKARI Space Telescopes, the infrared properties of galaxies have been explored at ever increasing sensitivity, spatial and spectral resolution.   The \\Spitzer Infrared Nearby Galaxies Survey (SINGS; Kennicutt et al. 2003) represents the `gold standard' study of nearby galaxies, employing every infrared instrument of \\Spitzer to study in detail the properties of 75 nearby `representative' galaxies.  A larger \\Spitzer imaging sample is found in the SINGS follow-up project, Local Volume Legacy (LVL). Expanding the sample to several thousand galaxies, the \\Spitzer Survey of Stellar Structure in Galaxies (S4G; Sheth et al. 2010) continued the SINGS and LVL surveys through the two short (near-infrared) wavelength bands of IRAC (3.6 and 4.5 $\\mu$m), focusing on the internal stellar structure of galaxies.   Following closely in succession to the AKARI all-sky survey (Murakami et al. 2006), the latest generation infrared space telescope, the Wide-field Infrared Survey Explorer (WISE),  expands these powerful surveys through its all-sky coverage and mega-pixel cameras, capable of constructing large, diverse and complete statistical samples--for both the near-infrared and mid-infrared windows--sensitive to both stellar structure (as with S4G) and interstellar processes (as with SINGS). WISE was specifically designed and optimized to detect and extract point source information. Detection, for example, was carried out using co-addition of image frames that were constructed with a resampling method based on a matched filter derived from the WISE point spread function (PSF).  As a consequence, this interpolation method tends to smear the images, making them less optimal for detection and characterization of resolved sources. However, owing to the stable PSF for all four WISE bands, there is a way to apply deconvolution techniques to recover from the smearing and further improve the spatial resolution. In this first paper of a two part series, we  demonstrate how the angular resolution of WISE may be enhanced to achieve information on physical scales comparable to those of \\Spitzer imaging, that enable detailed study of the internal anatomy of galaxies. We employ a resolution enhancement technique known as the Maximum Correlation Method (MCM;  Masci \\& Fowler 2009) to construct the WISE High Resolution Galaxy Atlas (WHRGA), consisting of several thousand nearby galaxies.  The WHRGA will comprise a complete mid-infrared source catalog and high-resolution image atlas of the largest angular-sized galaxies in the local universe. In this first paper we summarize the MCM algorithm and demonstrate its performance using WISE, \\Spitzer and GALEX imaging of nearby spiral galaxy NGC\\,1566. In the second paper (Jarrett et al. 2012b;  hereafter, referred to as Paper II), we demonstrate the early results of the WHRGA-project for a sample of 17 galaxies,  all observed by \\Spitzer and GALEX, chosen to be of large angular size, diverse morphology, and covering a range in color, stellar mass and star formation.   In addition to basic photometry, source characterization and surface brightness decomposition, Paper II also derives star formation rates and stellar masses.  This first paper is organized as follows.    Section 2 provides more technical information about the WISE mission and data products.  Section 3 the ancillary data used to compare to WISE, including \\Spitzer and GALEX imaging. Section 4 we outline the MCM image deconvolution- method and illustrate its performance using a set of simulated observations. Section 5 we focus on a case study of NGC\\,1566 to demonstrate the super-resolution performance using real WISE imaging data. All reported magnitudes are in the Vega System (unless otherwise specified). ",
        "conclusions": "In this paper we have presented a method by which to recover spatial information and significantly improve the angular resolution of WISE mid-infrared imaging, enabling detailed study of the internal anatomy of galaxies. Developed by the WISE Science Data Center, the Maximum Correlation Method (MCM) yields improvements that are 3 to 4 times better than the nominal WISE Atlas imaging resolution, and factors of 2 to 3 times improvement from standard `drizzle' co-addition.    Using the nearby galaxy NGC\\,1566 as a case study, we demonstrate how the angular resolution of WISE may be accurately enhanced to achieve information on physical scales comparable to those of \\Spitzer imaging. This method will be used to construct the WISE High Resolution Galaxy Atlas (WHRGA), consisting of several thousand nearby galaxies.  A pilot study is now underway, the initial results derived from a sample of 17 large, nearby galaxies is presented in a companion paper.     \\newpage   \\noindent Acknowledgements  This work is based [in part] on observations made with the {\\it Spitzer} and research using the  NASA/IPAC Extragalactic Database (NED) and IPAC Infrared Science Archive, all are operated by JPL, Caltech under a contract with the National Aeronautics and Space Administration. Support for this work was provided by NASA through an award issued by JPL/Caltech. R.J.A. was supported by an appointment to the NASA Postdoctoral Program at the Jet Propulsion Laboratory, administered by Oak Ridge Associated Universities through a contract with NASA. This publication makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration.  \\pagebreak"
    },
    "1208/1208.6014_arXiv.txt": {
        "abstract": "MESA (Modules for Experiments in Stellar Astrophysics) has become very popular among astrophysicists as a powerful and reliable code to simulate stellar evolution. Analyzing the output data thoroughly may, however, present some challenges and be rather time-consuming.  Here we describe MESAFace, a graphical and dynamical interface which provides an intuitive, efficient and quick way to analyze the MESA output. ",
        "introduction": "\\label{sec:Introduction}  The introduction of the new code for stellar evolution, MESA~\\cite{Paxton:2010ji} (Modules for Experiments in Stellar Astrophysics) represented somewhat of a revolution in the field of computational stellar evolution.  MESA is a one-dimensional stellar evolution code, organized in independent modules and continually updated. The code can evolve stars with a very wide range of initial masses and metallicities; it allows the specification of many parameters, for example, convection mechanism, mass loss, etc., and it can overcome difficulties typical of the previous stellar evolutionary codes. Notably, it can run a low mass star through the He-flash region.  All these aspects and the fact that the code is publicly available~\\cite{mesa} has made MESA extremely popular among astrophysicists and other physicists interested in stellar evolution and in how stars respond to new models of particle interactions.  Being publicly available, MESA is widely accessible and is continually tested by hundreds of scientists and students all over the world. In addition, several researchers have created new codes and routines to improve it and to simplify its use~\\cite{tools}.   A MESA run produces a very large amount of data on the life and structure of the star. This information is saved in several files: the file \"star.log\" includes global information about the star at different times whereas a large number of \\textit{profiles} (named \"log$n$.data\", where $n$ is an integer), contain detailed information about the profile of the star, each file referring to a specific age.  Though it is relatively easy, and for some purposes sufficient, to analyze the file \"star.log\", a complete analysis of the MESA output, which includes in some cases hundreds of  profiles, is much more challenging. In particular, to analyze the detailed structure of a star at a certain age, one should select the corresponding profile by reading the age of each of them (the information is provided in the $3^{\\rm nd}$ row of \"log$n$.data\"). Alternatively, one should read the \\textit{model number} corresponding to a certain age in the first column of the file \"star.log\", and then look for the log number associated with that age in the file \"profile.index\". Both ways are long and, in general, not very practical.  Another common problem is reading the age of a star from its position in the HR diagram (or any other plot which does not show the age explicitly, e.g., the \\textit{central temperature-central density} diagram). It would be useful to have a quick access to that information and, at the same time, to know which of the many profiles describes the structure of the star at that age.  Here we describe a graphical interface, MESAFace, which we created to address problems like the ones described above and, in general, to provide fast and simple access to \\textit{all the information} contained in the output of a MESA run, including all the profiles.  MESAFace is a very intuitive, efficient, and easy-to-use dynamical interface, which allows the selection of the information to be shown through buttons and menus and the age of the star to be selected through a slidebar.  The code is written in Mathematica and has been tested with Mathematica 7 and 8. However, a knowledge of Mathematica is not strictly necessary to use the interface.  MESAFace needs to access the MESA output data which, in a standard MESA run, are saved in the folder /work/LOGS. The data can also be downloaded directly from the web. Possibly soon, the results from some standard runs will be available from a dedicated  web space.  After the data has been loaded, MESAFace presents a graphical interface structured in three vertically organized panels: \\begin{itemize} \\item A history plot panel; \\item A profile plot panel; \\item An information panel. \\end{itemize} All the \\textit{controls} (buttons, menus, check-boxes, etc.) are contained on the left side of the interface. The user can decide what to plot and can change the age of the star by using these controls while the information shown on the right side is dynamically updated. Several aspects of the panels and graphs, for example, the color and style of the lines, can be customized by the user. A separate user manual~\\cite{userManual} explains in detail how to do that.    In this paper, we describe in detail the structure of the MESAFace code and how it accesses, manipulates, and presents the data.  The structure of the paper is the following: We begin with a brief summary of the MESA output, in section \\ref{sec:BriefDescriptionOfTheMESAOutput}. Later, in section \\ref{sec:TheStructureOfMf}, we give a basic description of the general structure of the MESAFace code. In section \\ref{sec:DownloadAndOrganizationOfTheFiles}, we describe with more details how MESAFace loads and organizes the files. In section \\ref{sec:TheInterface}, we discuss how it creates the dynamical interface and visualizes the data. Finally, in section \\ref{sec:Conclusions} we conclude with some overall comments.  A note on the style. We will use the \\texttt{Typewriter} style for the content of the Mathematica notebook, \"quotes\" for the name of the MESA files, and \\textbf{bold face} for the names of the controls in the interface. ",
        "conclusions": "\\label{sec:Conclusions}  MESAFace is a user-friendly and efficient interface to analyze the MESA output data. It is written in Mathematica, a very powerful and very well-documented mathematical software, and it is easily customizable.   The main purpose of this interface is to provide a tool to quickly analyze all the output data from a MESA run through graphs and tables. Our final objective was to have a very intuitive and easy-to-use instrument, accessible to researchers as well as to students. MESAFace allows the user to choose what to plot through various menus and to scroll the various profiles, organized by age, with a slidebar. The user can also easily zoom in different regions of the plot, change units, set a different zero for each axis, or plot the log (among other functions) of the data.  The production of high-quality graphs was not our priority, as other tools are available~\\cite{tioga}. However, improving the quality and create new options to manipulate the aspect of a graph, is a prominent direction for future improvement of MESAFace.  This paper is not meant to be a proper user manual. Basic instructions on how to use MESAFace are accessible on-line~\\cite{userManual}. Here, instead, we have described in detail the structure of the code and the meaning of the variables and controls. This allows the user not just to use, but if necessary, to modify and possibly improve this code."
    },
    "1208/1208.6364_arXiv.txt": {
        "abstract": "\\noindent In Basu et al. 2011 we reported the detection of Off-pulse emission from two long period pulsars B0525+21 and B2045--16. The pulsars were observed at a single epoch using the 325 MHz frequency band of the Giant Meterwave Radio Telescope (GMRT). In this paper we report a detailed study of the Off-pulse emission from these two pulsars using multiple observations at two different frequencies, 325 MHz and 610 MHz bands of GMRT. We report detection of Off-pulse emission during each observation and based on the scintillation effects and spectral index of Off-pulse emission we conclude a magnetospheric origin. The magnetospheric origin of Off-pulse emission gives rise to various interesting possibilities about its emission mechanism and raises questions about the structure of the magnetosphere. ",
        "introduction": "The basic pulsar phenomenon where a series of highly regular narrow pulses is observed has been explained from the very outset by invoking the lighthouse model. In this model the highly magnetized, rotating neutron star has a charge filled magnetosphere. The magnetic field surrounding the central star is characterized by a dipolar field with open and closed field line regions. The locus of the foot prints of the open field lines on the neutron star surface constitute the polar cap. It is believed that plasma is accelerated to relativistic velocities at the polar cap, which streams out along the open field lines giving rise to the observed coherent radio emission. This originate a few hundred kilometers above the stellar surface, and appear as narrow pulses due to the rotation of the neutron star. The plasma remains trapped in the closed field line regions thereby abstaining from the emission process.  The known cases of emission from pulsars outside the main pulse are the interpulse, pre/post-cursor emission and the pulsar wind nebulae (PWNe). The interpulse and pre/post-cursors emission appear as temporal structures in the time series data. The interpulse emission (located $180\\degr$ away from main pulse) are usually associated with orthogonal rotators, the radio emission in the interpulse believed to originate from the opposite magnetic pole. The pre/post-cursor emission are located closer to the main pulse and connected via a bridge emission. They are highly polarized and their location in the magnetosphere is uncertain. The time series observations of pulsars are insensitive to any temporally constant background emission. The technique of gated interferometry, where the main pulse emission is blocked, can probe a constant background emission by mapping the Off-pulse region. In the past gated interferometric studies have been employed to detect PWNe from very energetic young pulsars with spindown energies $\\dot{E} \\geq$ $10^{34}$ erg s$^{-1}$~\\citep{fra97} .  In recent studies we reported the detection of off-pulse emission from two long period pulsars based on gated interferometric observations with the Giant Metrewave Radio Telescope (GMRT) at 325 MHz~\\citep{bas11}, (hereafter PaperI). Our detection of off-pulse emission is a `unique' and surprising result because the two long period pulsars are relatively old and less energetic unlike the pulsars known to harbour PWN. Also the pulse profiles do not show any temporal structures outside the main pulse. We concluded on heuristic grounds that the detected Off-pulse emission is magnetospheric in origin based on outlandish expectations of ISM densities to sustain a PWM for these pulsar energetics.  In this paper we continue our previous investigations of Off-pulse emission. In section~\\ref{section2} we report the detection of Off-pulse emission at widely separated times and frequencies (see section~\\ref{subsec2.1}) and also account for the authenticity of these detections by carrying out a detailed test of the instrument to weed out possibility of the detected emission being a result of smearing of the pulsed signal across time (see section~\\ref{subsec2.2}). In section~\\ref{section3} we study the flux variations which lead us to an upper limit of the size of the Off-pulse emission region (see section~\\ref{subsec3.1}) and we also obtain secure estimates of its spectral index $\\alpha$ (see section \\ref{subsec3.2}). Based on our estimates of these properties we conclude a magnetopsheric origin for Off-pulse emission. Finally we discuss the results and their implications in section~\\ref{section4}.  \\begin{table*} \\begin{center} \\caption{ Observation details \\label{Obs_detail}} \\footnotesize{\\begin{tabular}{cccccccccc} \\tableline\\tableline Date & Pulsar & freq(chan) & Time & time res & Corr &         ang res       & On flux(rms) & Off flux(rms) &$\\frac{Off}{On}$ \\\\ &        &    MHz     & mins &   msec   &      & \\arcsec$\\times$\\arcsec&     mJy      &      mJy      &                 \\\\ \\tableline 18 Jan, 2010*& B0525+21& 325(128) & 160 & 261 &GHB& 9.5$\\times$6.5& 30.0$\\pm$2.1(0.55)&3.9$\\pm$0.5(0.45)&0.130$\\pm$0.026\\\\ 08 Jul, 2011 & B0525+21& 325( 32) & 240 & 250 &GSB& 8.3$\\times$6.7& 44.5$\\pm$2.4(0.65)&6.6$\\pm$0.5(0.4)&0.148$\\pm$0.019\\\\ 09 Jul, 2011 & B0525+21& 325( 32) & 120 & 250 &GSB&10.0$\\times$7.4& 48.2$\\pm$2.6(0.7)&6.5$\\pm$0.6(0.5)&0.135$\\pm$0.020\\\\ 15 Feb, 2011 & B0525+21& 610( 32) & 330 & 250 &GSB& 4.8$\\times$4.0& 21.1$\\pm$1.5(0.45)&3.6$\\pm$0.3(0.13)&0.171$\\pm$0.026\\\\ 22 Jul, 2011 & B0525+21& 610( 32) & 240 & 250 &GSB& 5.4$\\times$3.7& 20.0$\\pm$1.2(0.65)&2.8$\\pm$0.2(0.13)&0.140$\\pm$0.018\\\\ &               &     &     &   &               &             &           &               \\\\ 16 Jan, 2010*&B2045--16& 325(128) & 180 & 131 &GHB&11.9$\\times$7.2&110.5$\\pm$7.9(2.1)&4.3$\\pm$1.1(0.65)&0.039$\\pm$0.013\\\\ 03 Aug, 2011 &B2045--16& 325(256) & 200 & 125 &GSB&10.2$\\times$7.7& 57.5$\\pm$3.2(0.35)&6.8$\\pm$0.4(0.2)&0.118$\\pm$0.014\\\\ 14 Feb, 2011 &B2045--16& 610(128) & 180 & 131 &GHB& 5.8$\\times$4.3& 23.9$\\pm$1.7(1.0)&1.1$\\pm$0.2(0.13)&0.046$\\pm$0.012\\\\ 25 Aug, 2011 &B2045--16& 610(256) & 140 & 125 &GSB& 4.9$\\times$4.3& 43.2$\\pm$3.1(0.45)&4.6$\\pm$0.3(0.11)&0.106$\\pm$0.015\\\\   \\tableline\\tableline \\end{tabular}} \\tablenotetext{~}{\\small * reported in Basu et al. 2011} \\end{center} \\end{table*}  \\begin{figure} \\includegraphics[angle=0,scale=.5]{fig1.eps} \\caption{The folded profile (top) of B0525+21 representing the On- and Off-pulse gates. The corresponding contour maps of the On-pulse (bottom left) and Off-pulse (bottom right) emission from the pulsar. The data was recorded on 15 February, 2011 at 610 MHz. \\label{fig_P1imag}} \\end{figure}  \\begin{figure} \\includegraphics[angle=0,scale=.5]{fig2.eps} \\caption{The folded profile (top) of B2045-16 representing the On- and Off-pulse gates. The corresponding contour maps of the On-pulse (bottom left) and Off-pulse (bottom right) emission from the pulsar. The data was recorded on 25 August, 2011 at 610 MHz. \\label{fig_P2imag}} \\end{figure} ",
        "conclusions": "\\label{section4}  The discovery of radio emission in the Off-phase of two long period pulsars PSR B0525+21 and B2045--16 was reported in paperI. This emission, termed `Off-pulse emission', occur far away from the typical main pulse in the pulsar profile ($\\geq$ 80\\degr~from the peak of main pulse). The off-pulse emission unlike other emissions away from the main-pulse do not appear as temporal structures in the pulsar profile but rather as a constant background emission. \\footnote{A more sensitive phase resolved study in currently underway where we aim at imaging each phase resolved bin to localize the off-pulse emission region.}  In this paper we have focused on two important aspects:  (a) The off-pulse emission is unique in pulsars and it is important to ascertain that by no means the emission arises due to any instrumental effect or errors during data analysis. In section~\\ref{section2} we report detection of the off-pulse emission for B0525+21 and B2045--16 at widely separated times and different frequencies. We have conducted extensive tests and ruled out temporal leakage from a strong pulse as it passes through the GMRT receiver system as a possible source of Off-pulse emission. With these tests we are now absolutely certain that the off-pulse emission has an astronomical origin.  (b) The next important step was to establish through observational arguments the nature of Off-pulse emission. In section~\\ref{section3} we demonstrate that for PSR B0525+21 the Off-pulse to On-pulse flux ratio is constant at both 325 and 610 MHz, which in turn gives a spectral index $\\alpha_{OFF} \\sim$ --1.4. We have also demonstrated that despite apparent variations in the long term Off-pulse to On-pulse flux ratio for the pulsar B2045--16, it actually remains constant at all timescales when the Off-pulse flux variations are taken into account. This once again allowed us to determine the spectral index of $\\alpha_{OFF} \\sim$ --1.1. Based on the observed properties of refractive scintillation, we derived the radio emission region of the Off-pulse emission to have a maximum size of magnetospheric scale (see section~\\ref{subsec3.1}). The steep spectral index coupled with a highly compact emission region make it highly unlikely for the Off-pulse emission to be a PWNe.  The Off-pulse emission appear to be a completely new and hitherto undetected magnetospheric emission from pulsars. The estimated brightness temperature of the Off-pulse emission is greater than 10$^{18}$ K, assuming the emission to originate at the light cylinder. This strongly suggest a coherent radio emission process as the mechanism for explaining the Off-pulse emission which leads to the next important question as to where and how does this off-pulse coherent emission originate in the pulsar magnetosphere? Currently we do not have any good answers to these questions but there are a few likely possibilities that can be considered. The basic pulsar models suggest that the power from the rotational energy in a neutron star is tapped and converted into the observed radiation. The rotating magnetic field of the neutron star acts like a unipolar inductor, creating high electric fields around the neutron star where charged particles are accelerated to relativistic velocities along the open field lines. This eventually leads to the pulsar radiation. A large number of models exist which try to establish the location of the charge accelerating regions and suggest physical mechanisms that can excite the coherent radio emission in pulsars. However here we will not go into the details of any of these models, but would like to point out as why finding the location of the emission region in the pulsar magnetosphere is central in unravelling the origin of the Off-pulse emission.  In most pulsar theories the radio emission is excited due to development of plasma instabilities in the outflowing plasma. The main pulse emission arises around 500 km above the neutron star surface \\citep{bla91, ran93, kij98, krz09}. At these heights the magnetic field is very strong, and the relativistic charged plasma particles are forced to move along the curved magnetic field lines. Here the well known two-stream plasma instability naturally generates langmuir plasma waves. If the plasma is subjected to a non-stationary flow, then it leads to the modulational instability of Langmuir waves which results in formation of relativistic charged solitons capable of emitting coherent curvature radiation~\\citep{mel00, gil04, mit09}. The Off- pulse emission on the other hand lies in the region outside the main pulse, and currently we are not certain about its exact location. For the emission to originate from open field lines, the location needs to be above the main pulse in regions where the open dipolar magnetic field lines diverge. As the pulsar rotates, this divergent region above the main pulse is sampled by the observer as the Off-pulse emission. The other well known maser type plasma instability that can give rise to the coherent emission are the cyclotron-- cherenkov or cherenkov drift instability~\\citep{kaz87, kaz91}. The instabilities are due to the interaction of the fast particles of the primary beam with the normal modes of the electron-positron plasma. These instabilities are known to operate close to the light cylinder and could be a viable source of Off-pulse emission.  There might be other possible origin of the Off-pulse emission which we will need to evaluate and understand in the future, however it is clear that this newly found Off-pulse emission from low energetic slowly rotating neutron stars provide important clues in understanding the physical phenomenon operating in the pulsar magnetosphere."
    },
    "1208/1208.0297_arXiv.txt": {
        "abstract": "We propose a scenario that explains the apparent nitrogen deficiency in comets in a way consistent {\\bf with the fact that the surfaces of Pluto and Triton are dominated by nitrogen-rich ice}. We use a statistical thermodynamic model to investigate the composition of the successive multiple guest clathrates that may have formed during the cooling of the primordial nebula from the most abundant volatiles present in the gas phase. These clathrates agglomerated with the other ices (pure condensates or stoichiometric hydrates) and formed the building blocks of comets. We report that molecular nitrogen is a poor clathrate former, when we consider a plausible gas phase composition of the primordial nebula. This implies that its trapping into cometesimals requires a low disk temperature ($\\sim$20 K) in order to allow the formation of its pure condensate. We find that it is possible to explain the lack of molecular nitrogen in comets as a consequence of their postformation internal heating engendered by the decay of {\\bf short-lived radiogenic nuclides}. This scenario is found consistent with the presence of {\\bf nitrogen-rich ice covers on Pluto and Triton}. Our model predicts that comets should present xenon-to-water and krypton-to-water ratios close to solar xenon-to-oxygen and krypton-to-oxygen ratios, respectively. In contrast,  the argon-to-water ratio is predicted to be depleted by a factor of $\\sim$300 in comets compared to solar argon-to-oxygen, as a consequence {\\bf of poor trapping efficiency and radiogenic heating}. ",
        "introduction": "The abundance of the super volatile molecule N$_2$ in cometary nuclei is still under debate. Searches for UV fluorescence from N$_2$ with the Far Ultraviolet Spectroscopic Explorer (FUSE) have been unsuccessful (Bockel\\'ee-Morvan et al. 2004). Measurements of the N$_2$ abundance in 1P/Halley by Giotto were not feasible: owing to the low resolution of the mass spectrometer, it was not possible to discriminate N$_2$ from CO (Eberhardt et al. 1987). Ground-based observation of the N$_2^+$ band at 3914~\\AA~is a difficult challenge due both to the contamination of the $C_3$ (0,2,0)-(0,0,0) band and the presence of telluric N$_2^+$ emission lines. Different positive detections of this feature have been claimed in the ionic tail of comets Humason (1962 VIII) (Greenstein \\& Arpigny 1962), 1P/Halley (Lutz et al. 1993; Wyckoff \\& Theobald 1989), C/1987 P1 Bradfield (Lutz et al. 1993), and C/2002 VQ94 (LINEAR) (Korsun et al. 2006). Arpigny, from photographic spectra and other archival data estimated a positive N$_2^+$/CO$^+$ intensity ratio for 12 comets (Cochran et al. 2000). All these positive detections are based on low-resolution spectra. So far similar searches with high-resolution spectra have been unsuccessful. This was the case for observations of comets 122P/de Vico, C/1995 O1 (Hale-Bopp), and 153P/2002 C1 (Ikeya-Zhang) that did not detect N$_2^+$ band, yielding upper limits of $\\le$ 10$^{-5}$--10$^{-4}$ on the abundance of N$_2$ relative to CO (Cochran et al. 2000; Cochran 2002).  An interpretation of the possible N$_2$ deficiency in comets has been proposed by Iro et al. (2003). These authors assumed that comets were made of clathrate hydrates (hereafter clathrates) and argued that CO was preferentially incorporated in clathrates compared to N$_2$ after having investigated the competition between the trapping of these two molecules in clathrates. In order to explain the N$_2$ deficiency in comets, they argued that the nebula's temperature never cooled down below $\\sim$45 K in their formation region, impeding the formation of N$_2$-bearing ice (in clathrate or pure condensate form) that would have been agglomerated by comets. However, this mechanism is not consistent with the fact that Pluto and Triton {\\bf possess thick nitrogen ice covers} (Lellouch et al. 2011) while they are both expected to have been formed in the same region of the primitive nebula as ecliptic comets (Kavelaars et al. 2001). In addition, Iro et al. (2003) did not consider the competition between the trapping of N$_2$ in clathrates with that of the other abundant volatiles present in the nebula (CO, CO$_2$, CH$_4$, H$_2$S, ...). Indeed, due to their various propensities for trapping, the consideration of a larger set of molecules can drastically change the calculated fraction of N$_2$ incorporated in clathrates formed in the nebula.  Here we propose a scenario that could explain the apparent N$_2$ deficiency in comets in a way consistent with the presence of this molecule {\\bf on the surfaces of Pluto and Triton}. We use a statistical thermodynamic model to investigate the composition of the successive multiple guest clathrates (hereafter MG clathrates) that may have formed during the cooling of the primordial nebula from the most abundant volatiles present in the gas phase (see section 2). These clathrates agglomerated with the other ices and formed the building blocks of comets (hereafter cometesimals). The major ingredient of our model is the description of the guest-clathrate interaction by a spherically averaged Kihara potential with a nominal set of potential parameters. Our model allows us to report that N$_2$ is a poor clathrate former when considering a plausible gas phase composition of the primordial nebula, implying that its trapping into cometesimals requires a low disk's temperature ($\\sim$20 K) in order to allow the formation of its pure condensate. We find that it is possible to explain the lack of nitrogen in comets as a consequence of their postaccretion internal heating engendered by the decay of radiogenic nuclides (see section 3). % ",
        "conclusions": "{\\bf Our computations show that, under certain sets of conditions and evolutionary paths, initial nitrogen-rich cometesimals similar to Triton and Pluto may have evolved into nitrogen-depleted comets. The proposed scenario of selective trapping in clathrates in the nebula followed by outgassing due to short-lived radiogenic heating thus reconciles present-day observations with solar system formation models.} In our scenario, the composition of comets would have evolved as a result of radiogenic decay of short-lived nuclides. In the case of large objects like Pluto or Triton, the thermal evolution is more complex and involves partial or total physical differentiation. Although this significant processing could involve a substantial loss of volatiles, Pluto and Triton are actually large enough for most volatiles to be retained at the surface and in the atmosphere by gravity, as shown by Schaller \\& Brown (2007). Moreover, if accretion of these bodies were occurring in a cold environment and were slow--so that radiative removal of the impact heat were effective-- then Triton, Pluto and other large KBOs could have retained their volatiles and accreted fairly cold (see discussion by Canup \\& Ward (2002) in the context of the accretion of Ganymede and Callisto).  Our model predicts that comets should contain Xe/H$_2$O ($\\sim$3.6 $\\times$ 10$^{-7}$) and Kr/H$_2$O ($\\sim$5.0 $\\times$ 10$^{-6}$), ratios close to solar Xe/O and Kr/O, respectively. In contrast,  the Ar/H$_2$O ratio is predicted to be equal to $\\sim$1.7 $\\times$ 10$^{-5}$ in comets, a value about 300 times lower than solar Ar/O. In our model, the Ar depletion (compared to solar) in comets results from the poor propensity of this element to be trapped in different MG clathrates formed in the nebula. Because of this property, Ar was accreted by cometesimals only at low nebular temperature ($\\sim$20 K) and in pure condensate form. As a consequence, {\\bf the radiogenic heating above 22 K triggers the dissociation of clathrates and the outgassing of N$_2$ and sublimation of Ar.}  Measuring the noble gas abundances in comets might require in situ measurements via a high sensitivity mass spectrometer or ultraviolet spectroscopy. The ALICE ultraviolet spectrometer and the ROSINA instrument should put upper limits on the noble gas abundances in the coma of comet 67P/Churyumov-Gerasimenko (Balsiger et al. 2007; Stern et al. 2007) {\\bf that will probably help understand} the formation conditions of comets in the primordial nebula. However, the measurement of solar or subsolar noble gas abundances in comets will require a new generation of instruments such as the cryotrap, which is part of the mass spectrometer MASPEX developed in the context of the PRIME spacecraft project (Young et al. 2010)."
    },
    "1208/1208.0068_arXiv.txt": {
        "abstract": "In this work I study the problem of E/B-mode separation with binned cosmic shear two-point correlation function data.  Motivated by previous work on E/B-mode separation with shear two-point correlation functions and the practical considerations of data analysis, I consider E/B-mode estimators which are linear combinations of the binned shear correlation function data points.  I demonstrate that these estimators mix E- and B-modes generally. I then show how to define estimators which minimize this E/B-mode mixing and give practical recipes for their construction and use.  Using these optimal estimators, I demonstrate that the vector space composed of the binned shear correlation function data points can be decomposed into approximately ambiguous, E- and B-mode subspaces. With simple Fisher information estimates, I show that a non-trivial amount of information on typical cosmological parameters is contained in the ambiguous mode subspace computed in this formalism.  Next, I give two examples which apply these practical estimators and recipes to generic problems in cosmic shear data analysis: data compression and spatially locating B-mode contamination.  In particular, by using wavelet-like estimators with the shear correlation functions directly, one can pinpoint B-mode contamination to specific angular scales and extract information on its shape. Finally, I discuss how these estimators can be used as part of blinded or closed-box cosmic shear data analyses in order to assess and find B-mode contamination at high-precision while avoiding observer biases. ",
        "introduction": "\\label{sec:intro} \\setcounter{footnote}{0}  Cosmic shear, or the weak gravitational lensing of background galaxies by cosmological density fields, is one of the most important techniques for probing the properties of Dark Energy (see, e.g., \\citeauthor{weinberg2012} \\citeyear{weinberg2012} for a recent review) and also the growth of structure predicted by General Relativity (GR) or its possible modifications \\citep[e.g.,][]{schmidt2008,beynon2010,vanderveld2011}. Cosmic shear measurements can also help constrain other cosmologically interesting signals, such as primordial non-Gaussianity \\citep[e.g.,][]{fedeli2010,marian2011,maturi2011,giannantonio2012,hilbert2012}, the properties of various hot and warm dark matter models \\citep[e.g.,][]{schaefer2008,debono2010,markovic2011}, or the properties of neutrinos \\citep[e.g.,][]{cooray1999,song2004,hannestad2006,kitching2008,ichiki2009,debarnardis2009,jimenez2010}. To this end, ongoing and planned wide-field optical surveys, such as the DES\\footnote{The Dark Energy Survey - http://www.darkenergysurvey.org}, LSST\\footnote{Large Synoptic Survey Telescope - http://www.lsst.org}, Euclid\\footnote{http://sci.esa.int/euclid}, WFIRST\\footnote{Wide-Field Infrared Survey Telescope - http://wfirst.gsfc.nasa.gov}, HSC\\footnote{Hyper Suprime-Cam - http://www.naoj.org/Projects/HSC}, KIDS\\footnote{The Kilo Degree Survey - http://kids.strw.leidenuniv.nl}, and Pan-STARRS\\footnote{The Panoramic Survey Telescope \\& Rapid Response System - http://pan-starrs.ifa.hawaii.edu} surveys, will measure the shapes of hundreds of millions to billions of galaxies and thus cosmic shear signals with unprecedented statistical precision.  Given this incredible statistical power, understanding and mitigating potential systematic errors in these measurements will be very important.  Systematic contamination to cosmic shear signals can arise from a variety of sources, including the process of observing and estimating galaxy shapes from pixelated images \\citep[e.g.,][]{kaiser2000,bernstein2002,vale2004,hoekstra2004,guzik2005,paulinhenriksson2008,cypriano2010,voigt2010,kacprzak2012,refregier2012,voigt2012,antonik2012} or estimating photometric redshifts \\citep[e.g.,][]{ma2006,huterer2006,bridle2007,sun2009,hearin2010,bernstein2010,cunha2012}. There are also astrophysical sources of systematic errors, like intrinsic alignments \\citep[e.g.,][]{heavens2000,croft2000,catelan2001,crittenden2001,crittenden2002,jing2002b,lee2002,hirata2004,heymans2006b,hui2008,semboloni2008}, source galaxy clustering \\citep{schneider2002}, or the effects of baryons and galaxy formation on the matter power spectrum \\citep[e.g.,][]{white2004,zhan2004,huterer2005,jing2006,rudd2008,guillet2010,vandaalen2011,casarini2011,casarini2012,hearin2012}. These systematic errors, if left uncontrolled, can bias and/or degrade constraints on the properties of Dark Energy or modifications to GR from future surveys \\citep[e.g.,][]{hirata2003,hirata2004,guzik2005,huterer2005,huterer2006,mandelbaum2006b,ma2006,bridle2007,hirata2007,hearin2009,sun2009,semboloni2009,bernstein2010,hearin2010,kirk2011,kirk2012,cunha2012,hearin2012}.  Besides direct image and structure formation simulations to study cosmic shear data analysis, systematics, and theoretical modeling in detail (e.g., STEP1 \\citep{heymans2006}; STEP2 \\citep{massey2007}; GREAT08 \\citep{bridle2010}; GREAT10 \\citep{kitching2012}; \\citealp{jain2000,vale2003,lee2008,hilbert2009,sato2009,teyssier2009,hahn2010,kiessling2011,harnoisderaps2012}), it is important to distinguish between observational signals which can arise from GR and those which cannot \\citep{kaiser1992}.  At first order in the gravitational potential, GR will only produce cosmic shear patterns known as E-modes (see, e.g., \\citeauthor{dodelson2003} \\citeyear{dodelson2003} for a pedagogical introduction). The complementary patterns, know as B-modes, are not produced by GR at first order, though they can be produced in small amounts at higher order \\citep[e.g.,][]{jain2000,cooray2002,vale2003,hilbert2009,bernardeau2010,krause2010}.  Many of the sources of systematic contamination, though not all, can produce B-modes in addition to E-modes \\citep[e.g.,][]{crittenden2001,crittenden2002,schneider2002,vale2004,hirata2004,jarvis2004b,guzik2005,antonik2012}. Therefore assessing B-mode contamination in cosmic shear signals can test for systematic errors throughout the various steps of cosmic shear data analysis, from observing the galaxies with a telescope and imaging camera, all the way through to the theoretical modeling and constraints on cosmological parameters.  Methods for separating E- and B-modes in cosmic shear data have been studied extensively by previous authors. Broadly, these methods either operate directly on the shear field \\citep{schneider1998,seljak1998,hu2001b,heavens2003,leonard2012} or on the shear two-point correlation functions \\citep{crittenden2002,schneider2002,schneider2007,schneider2010,fu2010}. E/B-mode separation in the context of higher-order correlation functions has been studied as well \\citep{jarvis2004,schneider2005,shi2011,krause2012}. Additionally, techniques originally designed for the analysis of Cosmic Microwave Background polarization signals \\citep[e.g.,][]{wandelt2001,smith2006} can also be applied to cosmic shear \\citep{hikage2011}.  Importantly, the details of the implementation of these methods can effect their performance significantly \\citep[e.g.,][]{smith2006,kilbinger2006}. For the shear two-point correlation functions, $\\xi_{\\pm}(\\theta)$ defined below, \\citet{schneider2007} have shown that a broad class of E/B-mode statistics, can be written in the following form \\begin{eqnarray} E_{c} & = & \\frac{1}{2}\\int_{L}^{H}d\\theta\\,\\theta\\big[T_{+}(\\theta)\\xi_{+}(\\theta) + T_{-}(\\theta)\\xi_{-}(\\theta)\\big]\\nonumber\\\\ B_{c} & = & \\frac{1}{2}\\int_{L}^{H}d\\theta\\,\\theta\\big[T_{+}(\\theta)\\xi_{+}(\\theta) - T_{-}(\\theta)\\xi_{-}(\\theta)\\big]\\nonumber\\ . \\end{eqnarray} By choosing the range of integration $[L,H]$ and the forms of $T_{\\pm}(\\theta)$ properly, one can show that $E_{c}$ will contain only E-mode information and $B_{c}$ will contain only B-mode information either over an infinite interval or finite interval with $L >0$ \\citep{schneider2007}.  Note that these statistics assume one has continuous shear correlation function data.  Practical implementations of the $E_c$ and $B_c$ statistics in cosmic shear data analysis are constrained in several ways.  The shear correlation functions are usually estimated in bins of angle, say $N$ bins, with some effective binning weight $W_{i}(\\theta)$.  In particular, the expectation value of the estimated shear correlation function data point for the $i$th bin is \\citep[e.g.,][]{schmidt2009a} \\begin{equation}\\label{eqn:xpmdef} \\left<\\widehat{\\xi}_{\\pm i}\\right> = \\int_{L_{i}}^{H_{i}}d\\theta\\,W_{i}(\\theta)\\xi_{\\pm}(\\theta)\\ . \\end{equation} The form of these binning functions and this result is discussed below and in Appendix~\\ref{app:srcclust}.  Thus the statistics $E_{c}$ and $B_{c}$ in this case must be estimated from the binned shear correlation function data, $\\widehat{\\xi}_{\\pm i}$.  Also, in order for the statistics $E_{c}$ and $B_{c}$ to remain pure two-point statistics, any procedure for estimating them from cosmic shear data can only consider linear combinations of the binned shear correlation function data, \\begin{equation}\\label{eqn:xpmlinstat} X_{\\pm}=\\frac{1}{2}\\sum_{i}\\left[F_{+i}\\widehat{\\xi}_{+i} \\pm F_{-i}\\widehat{\\xi}_{-i}\\right]\\ , \\end{equation} where the $F_{\\pm i}$ are constants which describe the statistics \\citep[see e.g.,][]{kilbinger2006b}.  In this work, I study the use of these linear combinations for E/B-mode separation and thus the effects of these constraints on cosmic shear data analysis. In particular, after covering the basics of cosmic shear in Section~\\ref{sec:cosmodefs}, I demonstrate in Section~\\ref{sec:ebmix} that \\textit{even if the statistics $E_{c}$ and $B_{c}$ contain pure E- and B-mode information, the statistics $X_{\\pm}$ generally exhibit E/B-mode mixing due to the binning}.  In Section~\\ref{sec:buildEBest}, I show how to define the statistics $X_{\\pm}$ through the choice of the $F_{\\pm i}$, so that they suppress the E/B-mode mixing below detectable levels for any current or upcoming cosmic shear survey.  In this section I also give practical recipes for computing the $F_{\\pm i}$.  Other potential and ultimately equivalent optimal estimator definitions are discussed in Section~\\ref{sec:otherest}.  I then show how to use these optimal estimators to divide the vector space of correlation function data points up into approximately ambiguous, E- and B-mode subspaces in Section~\\ref{sec:ebmodedec}.  I compute the Gaussian covariances of these statistics in Section~\\ref{sec:ebcov}.  In Sections~\\ref{sec:cosebi} and \\ref{sec:wvstats}, I provide two examples of these statistics to illustrate how to build and use them in practice.  I provide an example of how to decompose the shear correlation functions into ambiguous, E- and B-modes and also discuss the Fisher information content of these subspaces in Section~\\ref{sec:fishinfo}. I find that the ambiguous mode subspace has a non-trivial amount of information about typical cosmological parameters.  Finally, I conclude and discuss how these statistics are applicable to blinded or closed-box cosmic shear data analyses in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} In this work, I have demonstrated that the use of two-point E/B-mode estimators with binned shear correlation function data will generally result in non-trivial E/B-mode mixing.  I computed the amount of this mixing and provided new E/B-mode estimators which minimize the unwanted mode mixing. I also gave practical recipes for building and using these estimators with binned cosmic shear data.  Using these optimal estimators, I demonstrated that approximate decompositions of the vector space of binned correlation function points into ambiguous, E- and B-modes do exist.  Using one of these decompositions, I found that the ambiguous mode subspace contains a non-trivial amount of information on typical cosmological parameters.  I also gave two example applications of these new estimators to generic problems in cosmic shear data analysis, data compression (Section~\\ref{sec:cosebi}) and B-mode localization with wavelets (Section~\\ref{sec:wvstats}).  The estimators presented here have several nice features adapted to practical cosmic shear data analysis. \\begin{itemize} \\item They are linear combinations of the binned shear correlation function data, defined in Equation~(\\ref{eqn:xpmlinstat}), and thus treat the binning explicitly.  This property makes their computation and also error propagation/analysis with them trivial once the shear correlation function and its errors are known. \\item The level of E/B-mode mixing due to the binning can be computed exactly, up to the knowledge of the binning window functions, for these estimators using Equations~(\\ref{eqn:xipmell}) and (\\ref{eqn:wpmdef}).  In the limit of small bins, which is needed to suppress the E/B-mode mixing, the binning window functions are expected to be close to the geometric approximation used throughout this work. \\item They give quantitative criteria in terms of the E/B-mode mixing by which to decide the number of shear correlation function bins to use, demonstrated in Figure~\\ref{fig:cosebiss2n} for the COSEBI-like estimators. \\item The design of new E/B-mode estimators with specific purposes using the formalism presented in this work reduces to simple linear algebraic manipulations, presented in Section~\\ref{sec:buildEBest}. \\end{itemize}  The optimal statistics presented in this work are well-suited to blinded or closed-box, high-precision cosmic shear data analyses.  For example, before any shear correlation functions are computed from the data, one can estimate given the expected statistical accuracy of the data, the exact amount of E/B-mode mixing for a given binning scheme and set of estimators. One can then choose a fiducial binning scheme and estimator choice that properly minimizes the E/B-mode mixing and retains all of the cosmological information.  Then these choices can be fixed throughout the data analysis process in order to avoid any observer biases in detecting or assessing B-mode contamination arising from how exactly the E/B-mode separation was done. Additionally, in a blinded or close-box analysis one might not look at plots of the shear correlation function data until the entire analysis is complete. Unfortunately in this case, one might miss crucial information about potential systematics in the \\textit{shape} of the shear correlation function data. The wavelet-like B-mode estimators presented in this work can be used as a substitute for and quantitative measure of the information gained by looking at the shape of the shear correlation function data.  Additionally, they can be applied in automated way to large data sets in order to pinpoint areas of potential systematic contamination for further investigation.  Importantly, because these estimators have a large degree of spatial locality, they can potentially provide crucial information on where any B-mode contamination is coming from and not just indicate its existence.  The ability to easily and quickly design E/B-mode estimators tailored to a specific purpose can potentially be very useful in practice as well.  For instance, the problem of computing an E/B-mode statistic which maximizes the signal-to-noise or cosmological information content, as considered by \\citet{fu2010}, could now be reformulated with the linear estimators presented here.  Additionally, one could build estimators which are along the normal modes of the correlation function data computed from the correlation function covariance matrix.  Also, one could attempt to build direct estimators for the E/B-mode correlation functions \\citep{crittenden2002,schneider2002,schneider2010} over a finite interval.  Alternatively, one could attempt to build filters which are localized in Fourier space in order to exclude certain wave modes, for example to mitigate the effects of uncertainty in the matter power spectrum at small scales \\citep[see e.g.,][]{huterer2005b}, or to be sensitive to B-modes from only a range of $\\ell$. As simple linear combinations of cosmic shear data points, the estimators presented in this work are well-suited to these applications and to practical cosmic shear data analysis in general."
    },
    "1208/1208.0404_arXiv.txt": {
        "abstract": "We report the existence of spiral arms in the recently formed gaseous and dusty disk of the closest giant elliptical, NGC~5128 (Centaurus A), using high resolution $^{12}$CO(2--1) observations of the central 3\\arcmin\\ (3~kpc) obtained with the Submillimeter Array (SMA). This provides evidence that spiral-like features can develop within ellipticals if enough cold gas exists. We elucidate the distribution and kinematics of the molecular gas in this region with a resolution of 4\\farcs4 $\\times$ 1\\farcs9 (80~pc $\\times$ 40~pc). The spiral arms extend from the circumnuclear gas at a radius of 200~pc to at least 1~kiloparsec. The general properties of the arms are similar to those in spiral galaxies: they are trailing, their width is $\\sim$ 500 $\\pm$ 200~pc, and the pitch angle is $\\sim$20$\\rm ^o$. From independent estimates of the time when the HI-rich galaxy merger occurred, we infer that the formation of spiral arms occured on a time scale of less than $\\sim$10$^8$~yr. The formation of spiral arms increases the gas density and thus the star formation efficiency in the early stages of the formation of a disk. ",
        "introduction": "\\label{sec1}   Some ellipticals show rotating disk-like features of gas and dust, expected to form as a result of gas accretion from the intergalactic medium, cannibalization of other galaxies, or re-accretion of material originally expelled during the major merger event that resulted in the formation of the elliptical itself \\citep[e.g.]{1997SSRv...81....1H,2002AJ....124..788Y,2011MNRAS.417..882D,2011MNRAS.414..940Y}.  % \\citet{1993ApJ...419..544S} proposed an evolutionary sequence of disk formation in dust lane ellipticals via accretion, from irregularly distributed gas to more evolved and settled disks. However, our understanding of the properties and evolution of such relatively settled disks is poorly understood yet, partly because most observations of the gas in elliptical galaxies are hampered by poor angular resolution and/or sensitivity. In particular, it is not known if these newly created disks in elliptical galaxies are uniform once settled, or whether deviations from axisymmetry such as prominent spiral structures can develop, affecting the subsequent star formation patterns, and the overall evolution of the disk itself.    At a distance of D $\\simeq$ 3.8~Mpc \\citep{2010PASA...27..457H} (where 1\\arcsec\\ roughly corresponds to 18~pc), NGC~5128 (hereafter Cen~A) is the closest giant elliptical \\citep{1976ApJ...208..673V,1998AARv...8..237I,2012AJ....143...84H} with a prominent dust lane along its minor axis \\citep{1998AARv...8..237I,2012AJ....143...84H} (see Figure~\\ref{fig1}, left). This dust lane contains large amounts of gas and dust, as traced by H$\\alpha$ \\citep{1992ApJ...387..503N}, mid-IR \\citep{2006ApJ...645.1092Q}, \\ion{H}{1} \\citep{1990AJ.....99.1781V,2010AA...515A..67S}, and CO \\citep{1992ApJ...391..121Q,1993AA...270L..13R,2009ApJ...695..116E}.  The \\ion{H}{1} emission still shows unsettled gas at r$>$6~kpc, including tail/arm like structures, and it is estimated that the merger event occurred within a timescale of only 0.3~Gyr  \\citep{2010AA...515A..67S}. The gas in molecular phase and the dusty component inside the inner few kiloparsecs are well settled within a warped disk of about 7\\arcmin\\ with an inclination of $i$ $\\sim$ 70$^{\\rm o}$ \\citep{2006ApJ...645.1092Q}, although its detailed distribution at scales of a few hundreds of parsec has hitherto been poorly understood.  The most popular model to reproduce the large scale disk structure is that of an homogenous warped and thin disk (e.g. \\citealt{1992ApJ...387..503N,1992ApJ...391..121Q,2006ApJ...645.1092Q,2009AA...502L...5K,2010AA...515A..67S,2010PASA...27..396Q}). According to these models, the disk crosses the line of sight at two galactocentric radii. One of the most remarkable features in favor of this model is that in projection it can reproduce the parallelogram structure seen in the IR, having a major side of $\\sim$3\\arcmin\\ (3~kpc) along a PA = 120$\\rm ^o$ \\citep{2006ApJ...645.1092Q}.  However,  the inner 1 kiloparsec of Cen~A (up to a few 100 pc from the nucleus) is very complex. \\citet[][hereafter Paper I]{2009ApJ...695..116E} report $^{12}$CO(2--1) SMA observations of its nuclear region in a single pointing, with a resolution of 6\\farcs0 $\\times$ 2\\farcs4 (100 $\\times$ 40 pc), revealing the detailed distribution and kinematics of the molecular gas in the central inner kpc, including a compact circumnuclear disk of molecular gas surrounding the AGN, as well as outer molecular gas partly associated with the parallelogram structure. While the warped disk models reproduce the observed features (distribution, kinematics, dust lane appearance) at a large scale, Paper I shows that the molecular gas distribution in the inner kpc is not sufficiently in agreement with these models. A weak bi-symmetric potential was proposed in Paper I to explain the deviations with respect to this model.  We present new observations using the Submillimeter Array (SMA\\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics, and is funded by the Smithsonian Institution and the Academia Sinica.}; \\citealt{2004ApJ...616L...1H}), that allow us to shed more light into the complex few inner kiloparsecs. The resolution of the mosaic we present here is a factor of 45 higher than existing $^{12}$CO(2--1) maps covering the dust lane, with a resolution of $23\\arcsec$ (or $\\sim380$~pc) \\citep{1993AA...270L..13R}. Our mosaic covers an area three times larger than previously published interferometric CO maps (Paper I) to fully cover the entire parallelogram structure. ",
        "conclusions": "A consequence of the compression produced by these spiral arms is that it is expected to trigger star formation (SF) on the leading edge. This is consistent with abundant SF traced by Pa$\\alpha$ \\citep{2000ApJ...528..276M} for example, coincident at least with the southern molecular spiral arm (nearest part and thus likely the least obscured). The spiral features seen in the $^{12}$CO emission is also associated with the distribution of certain tracers observed in poorer angular/spectral resolution maps, such as the SCUBA 450$\\mu$m emission map in \\citet{2002ApJ...565..131L}, as well as the pure rotational line of molecular hydrogen H$_2$ (J=2--0) S(0) ($28.22~\\mu$m) emission observed with Spitzer/IRS by \\citet{2008MNRAS.384.1469Q}. The molecular hydrogen transition H$_2$ (J=2--0) S(0)  indicates the presence of gas with T $\\sim$ 200~K \\citep{2008MNRAS.384.1469Q}, which is likely tracing photodissociation regions associated with abundant star formation.  A small perturbation could have triggered this spirality, such as a non-axisymmetric weak potential (Paper I) or a minor merger after the disk was relatively well settled \\citep[e.g.][]{2011Natur.477..301P}. From the \\ion{H}{1} structure and kinematics we infer that the formation of spiral arms took place in less than 0.3~Gyr \\citep{2010AA...515A..67S}, which is likely the most accurate measure of time at our disposal since \\ion{H}{1} is one of the most sensitive components of interaction. Although evidence suggests that spiral arms are transient \\citep{2011MNRAS.410.1637S,2011ApJ...735..101F}, recent simulations claim that spiral structures may be long lived \\citep{2012arXiv1204.0513D}. 3-D self-gravitating models for late type spirals suggest that the pitch angle we find here would correspond to a long-lasting feature, because the density response tends to avoid larger pitch angles \\citep{2012ApJ...745L..14P}. Similar simulations in the deep potentials of giant ellipticals would be needed to investigate the response of the gas.  The study of the properties of recently accreted molecular gas in the deep potential wells of elliptical galaxies is a powerful tool in disentangling how galaxies form and evolve. Recent smoothed-particle hydrodynamic (SPH) cosmological simulations in the $\\Lambda$ Cold Dark Matter (CDM) scenario \\citep{2011ApJS..192...18K} are able to reproduce most of the properties of disks \\citep{2011MNRAS.410.1391A,2011ApJ...742...76G,2012MNRAS.tmp.2503D} at the current epoch,  including their fine structure morphology, the observed angular momentum, the Tully-Fisher relation, and the SFR to gas surface density ratio (Kennicutt-Schmidt). However, to understand galaxy evolution it is essential to anchor these simulations with the actual properties of nascent disk galaxies. The properties of disks, in particular the existence of spiral arms, in these systems which are in the early stages following significant accretion can be compared with numerical simulations. Not only can we compare the observed gas distribution and kinematics of recently formed disks to models, but also the effect of SF and Active Galactic Nuclei (AGN) feedback in the disk evolution. Moreover, with the advent of ALMA, it is now possible to validate theories of disk formation as they enter a deep gravitational potential, performing similar studies in objects along the evolutionary sequence of accreted gas in dust lane ellipticals \\citep{1993ApJ...419..544S},  from systems whose gas is still irregularly distributed to more evolved and settled disks such as in Cen~A."
    },
    "1208/1208.5690_arXiv.txt": {
        "abstract": "During the last taking data campaigns of the CAST experiment, the micromegas detectors have achieved background levels of $\\approx 5 \\times 10^{-6}$keV$^{-1}$cm$^{-2}$s$^{-1}$ between 2 and 9 keV. This performance has been possible thanks to the introduction of the microbulk technology, the implementation of a shielding and the development of discrimination algorithms. It has motivated new studies towards a deeper understanding of CAST detectors background. One of the working lines includes the construction of a replica of the set-up used in CAST by micromegas detectors and its installation in the Canfranc Underground Laboratory. Thanks to the comparison between the performance of the detectors underground and at surface, shielding upgrades, etc, different contributions to the detectors background have been evaluated. In particular, an upper limit $< 2 \\times 10^{-7}$keV$^{-1}$cm$^{-2}$s$^{-1}$ for the intrinsic background of the detector has been obtained. This work means a first evaluation of the potential of the newest micromegas technology in an underground laboratory, the most suitable environment for Rare Event Searches. ",
        "introduction": "\\subsection{CAST}  \\medskip The principle of an axion-helioscope\\cite{Sikivie} consists of the conversion into photons of some of the axions\\cite{axions1}\\cite{axions2} generated and comming from the Sun in presence of a transverse magnetic field. The solar axion flux could then be identified as an excess of photons registered by X-ray detectors (axion's spectrum is ranges from 2 to 9 keV) during the time the helioscope is aligned with the Sun.  \\medskip The CERN Axion Solar Telescope, CAST\\cite{cast}\\cite{exrs}, is the best realization until now of an axion-helioscope thanks to the powerful LHC dipole prototype and to the sensitivity of its X-rays detectors. The typical background rate during current operation phase in CAST\\cite{newcast} has been $5 \\times 10^{-6}$keV$^{-1}$cm$^{-2}$s$^{-1}$ from 2 to 9 keV. The goal of improving this performance towards a new generation of axion helioscopes\\cite{NGAH}, needs a carefully characterization of present detectors.  \\subsection{CAST micromegas detectors}  \\medskip  There are three main requirements for low background detectors: the use of low radicative materials, a shielding againts environmental radiation and good discrimination capabilities. The successive introduction and continuous improvement of these strategies has led to a constant decrease of the background level in CAST micromegas detectors.  \\medskip In the new generation of micromegas\\cite{micromegas}, called microbulk\\cite{microbulk}\\cite{Paco_microbulk}, all the readout is contained in a 80 $\\mu$m foil. This foil has been measured showing good radiopurity\\cite{paquito}. The rest of the mechanical structure of the detector is made of Plexiglass (also radiopure) with the exception of the drift cathode which is made of steel. A thin window (only 4 $\\mu$m of aluminized mylar) and the operation at slighty overpressure (1.35 bar) optimizes the detector efficiency for X-rays detection.  \\medskip The detector is coverd by a shielding composed of a first 5 mm thickness copper layer and second one of archeological lead and 25 mm thickness. Also N$_2$ flow is used to avoid Rn emanations nearby the detector. Moreover the readout design allows a descriptive characterization of the registered events, providing both spatial and time features. This information is used to discriminate X-rays events from other interactions\\cite{analysis}\\cite[Chap. 4 and 5]{Javis}.  \\begin{figure}[!h] \\centering \\includegraphics[width = .95\\textwidth, angle = 0]{CASTdetectors.png} \\caption{From left to right: pixelized strip readout; micromegas foil (containing amplification structure plus readout plane) glued on a PCB racket; addition of the chamber with a thin aluminized window and the electronic cards; the detector connection to the magnet pipe during an intermediate installation where lead shielding can be observed.} \\label{fig:castdetectors} \\end{figure}   \\subsection{Work motivation and methodology}  A new set-up (Figure 2), that essentially reproduces the CAST one (as described in the previous section), was built in the University of Zaragoza. It is completed with a CAST microbulk detector; the experiments acquisition is also replicated and the same analysis programs are used to process the data. The set-up was equiped with a calibrator containing a $^{55}$Fe source in order to check periodically the detector status.  \\begin{figure}[!h] \\centering \\includegraphics[width = .48\\textwidth]{LSC-setup.png} \\includegraphics[width = .48\\textwidth]{DSC_0054little.JPG} \\caption{Left: CAST-like test set-up first used in Zaragoza's test bench and later in LSC. Light shielding composed by 2.5 cm of lead with an internal shielding of 0.5 cm of more radiopure copper. The two last lead caps are added after closing the Faraday cage. The set-up is prepared for 6 keV X-Rays automatic calibrations thanks to a $^{55}$Fe source inside the Faraday cage. Right: the same set-up after the later upgrading: a 20 cm thick lead shielding. The liquid nitrogen dewar and the acquisition electronics are shown too.} \\label{fig:LSCset-up} \\end{figure}  \\medskip The goal of the work is to study the influence of changes in the detector environment on the detector background. For each change a concrete origin of background is pretended to be avoided or substantially reduced. Comparisons between different set-ups and detectors were made too, trying to itemize the several causes which are susceptible to have an influence and evaluate its strength.  \\medskip Moving the detector to the Canfranc Underground Laboratory, at 2500 m.w.e depth, the cosmic muons are reduced by a factor $\\times 10^{4}$\\cite{LSCmuons}. In contrast, the radon concentration is much more intense than at surface and highly variable, with typical values from 100 to 200 Bq/m$^3$. The absence of cosmic rays allows the use of a much thicker shielding, which practically blocks the environmental radiation. ",
        "conclusions": "\\medskip LSC is a useful test bench for CAST detectors as it allows us to analyze the relative weight of the different physical contributions to CAST micromegas background. A direct comparison of the background by using the same set-up at surface and underground shows that, even when cosmic rays dominate the trigger rate at surface, muons are efficiently rejected by the micromegas. An important part of the CAST background, in special the fluorescence peaks, is related with the intrusion of gammas from the connection between the detector and the magnet bores. This is already being reduced by means of upgrading the shielding of the line.  \\medskip CAST micromegas final background level is dominated by the environmental gamma flux, as it is clear from the reduction obtained by shielding upgrade: about 30 times, which leads to an upper limit for the micromegas detector intrinsic background below $2 \\times 10^{-7}$keV$^{-1}$cm$^{-2}$s$^{-1}$ in the 2-9 keV range. The work is ongoing to identify more concrete origins for the internal contaminations.  \\begin{figure}[!h] \\centering \\includegraphics[width = .43\\textwidth, angle = 0]{Bkg_History.png} \\caption{Historical background levels for CAST micromegas detectors.} \\label{fig:history} \\end{figure}  \\medskip From this work a new point is added to the CAST micromegas background historical curve (Figure 5). The new point goes away the general tendency, roughly exponential, because it does not correspond to the context of the true CAST experiment. It must be considered as a limit for the future evolution of CAST micromegas background imposed by the intrinsic radioactivity. Although new techniques may bring more radiopure micromegas. It can also represent a first indication for micromegas in a new context: the Underground Physics."
    },
    "1208/1208.2286_arXiv.txt": {
        "abstract": "We introduce a topological flux function to quantify the topology of magnetic braids: non-zero, line-tied magnetic fields whose field lines all connect between two boundaries. This scalar function is an ideal invariant defined on a cross-section of the magnetic field, and measures the average poloidal magnetic flux around any given field line, or the average pairwise crossing number between a given field line and all others. Moreover, its integral over the cross-section yields the relative magnetic helicity. Using the fact that the flux function is also an action in the Hamiltonian formulation of the field line equations, we prove that it uniquely characterizes the field line mapping and hence the magnetic topology. ",
        "introduction": "In many plasmas, ranging from astrophysics to magnetic confinement fusion, the topology---i.e., the linking and connectivity of the magnetic field lines \\cite{longcope2005}---is an approximate invariant of the dynamics. This is because these plasmas typically have such low dissipation that, to first approximation, their evolution is ideal. That is, on large scales where the magnetohydrodynamic approximation holds, they satisfy an ideal Ohm's law \\cite{goedbloed2004}, preserving the magnetic topology. Therefore a practical question is, given two magnetic fields satisfying the same conditions on the boundary of some volume $V$, can one be reached from the other by some ideal evolution in $V$?  We restrict our attention to line-tied magnetic flux tubes, where all field lines stretch between two boundaries and the magnetic field in the volume is non-vanishing. This models, for example, a coronal loop in the Sun's atmosphere, where the footpoints remain essentially fixed on the rapid timescale of coronal relaxation \\citep[e.g.,][]{parker1972,*vanballegooijen1985,*mikic1989,*longcope1994,*galsgaard1996, *browning2008,pontin2010}. To simplify the presentation in this paper, we consider a magnetic field ${\\bf B}$ defined on a cylinder \\mbox{$V=\\{(r,\\phi,z) | 0\\leq r\\leq R, 0\\leq z\\leq 1 \\}$}, satisfying $B_z>0$ everywhere in $V$, and impose the boundary conditions that ${\\bf B}|_{\\partial V}={\\bf e}_z$ and ${\\bf v}|_{\\partial V}=0$, where ${\\bf v}$ is the velocity. Extensions of the results to more general boundary conditions are discussed in Section \\ref{sec:conclusion}. For convenience, we call magnetic fields satisfying the above conditions ``magnetic braids'' (Fig. 1). Two magnetic braids are topologically equivalent if they can be linked by an ideal evolution where ${\\bf v}=0$ on $\\partial V$ throughout.  In principle, one can determine whether two magnetic braids are topologically equivalent by comparing their field line mappings from $z=0$ to $z=1$. Throughout this paper, let $f(x_0; z)\\in V$ denote the point at height $z$ on the field line traced from $x_0\\equiv(r_0,\\phi_0,0)$ on the $z=0$ boundary. Under this parameterization of field lines by $z$, we have \\begin{equation} \\frac{\\mathrm{d} f(x_0; z)}{\\mathrm{d}z} = \\frac{{\\bf B}\\big(f(x_0; z) \\big)}{B_z\\big(f(x_0; z) \\big)}. \\end{equation} For shorthand, we shall denote the mapping from $z=0$ to $z=1$ as $F(x_0)\\equiv f(x_0;1)$. Under our boundary conditions, two magnetic braids ${\\bf B}$, $\\widetilde{\\bf B}$ are equivalent if and only if $F=\\widetilde{F}$. Note that if we were to relax the condition that ${\\bf B}={\\bf e}_z$ on the side boundary $r=R$, then $F$ would determine the topology only up to an overall rigid rotation through $2n\\pi$, $n\\in\\mathbb{Z}$. Mathematically, field line mappings are symplectic, since they preserve magnetic flux. Symplectic mappings have long been used themselves as models of periodic magnetic fields in fusion devices \\cite{morrison2000,*caldas2012}, and field line mappings have also been used extensively for characterizing line-tied coronal magnetic fields \\cite{titov2002,*titov2007,*yeates2012,*pariat2012}. But the mapping is usually very sensitive to small fluctuations in the underlying magnetic field. This makes it very difficult, if not impossible, to determine whether two field line mappings can be related by an ideal evolution in anything other than highly idealized situations.  A much more robust topological quantity is the total magnetic helicity, which has a broad range of applications in both laboratory and astrophysical plasmas \\cite{brown1999}. It has proved so robust that it has been hypothesized to be the only quantity determining the final state of turbulent relaxation in reversed-field pinches and similar devices \\cite{taylor1974,*taylor2000}. But the helicity is an extreme reduction of the topological information in the three-dimensional magnetic field to a single number, and it does not uniquely characterize the topology. There are a large class of magnetic fields that have the same helicity but different field line mappings.  In this paper, we describe a quantitative measure ${\\cal A}(x_0)$ which we call the ``topological flux function''. This is a scalar function defined on a cross section of the magnetic field, or equivalently on each field line. It is more robust than the field line mapping, while containing more detailed information than the helicity. A similar function was introduced for magnetic fields in a half-space by Berger \\cite{berger1988}, who showed that it is effectively a helicity per field line. The function ${\\cal A}$ has also appeared in the literature in a different guise: as an action integral yielding the magnetic field line equations in a variational formalism \\cite{cary1983,morrison2000,petrisor2002,*apte2006}. These interpretations are discussed in more detail in Section \\ref{sec:interp}. Recently, we have described how ${\\cal A}$ may also be viewed as a generalization of the scalar flux function used to define two-dimensional magnetic fields, and how it may be used to define and measure magnetic reconnection \\cite{yeates2011}. Here we go further. Our main result is that, with a particular choice in  its definition, ${\\cal A}$ uniquely characterizes the topology of a magnetic braid. In other words, ${\\cal A}=\\widetilde{\\cal A}$ for two magnetic braids if \\emph{and only if} they are topologically equivalent. Moreover, not only can ${\\cal A}$ determine whether two braids are topologically equivalent, it can also quantify how much dissipation or reconnection is needed to connect the two states. This will be invaluable in future dynamical studies of such systems.  This paper is organized as follows. In Section \\ref{sec:def}, we define ${\\cal A}$ and give its basic physical interpretation in terms of magnetic flux. Its interpretations as a field line helicity, as an average crossing number, and as a Hamiltonian action are described in Section \\ref{sec:interp}. In Section \\ref{sec:main} the Hamiltonian interpretation is used to prove our main result that ${\\cal A}$ uniquely characterizes the magnetic topology. We outline in Section \\ref{sec:conclusion} how the boundary conditions may be relaxed. ",
        "conclusions": "\\label{sec:conclusion}  For clarity of presentation, we have assumed certain boundary conditions on the magnetic field, namely that ${\\bf B}=1$ on all boundaries of our cylinder $V$, and that ${\\bf v}=0$ on the bottom and top boundaries. We indicate here how the results generalize when these conditions are relaxed.  If one allows $B_\\phi\\neq 0$ on the side boundary $r=R$, then $G$ is determined only up to an overall rigid rotation, and we would have the result that $\\widetilde{\\cal A}$ and ${\\cal A}$ differ by a constant if and only if $\\widetilde{F}$ and $F$ differ by an overall rigid rotation. Such a rotation can be detected from knowledge of $F$ and $\\widetilde{F}$ on the side boundary alone. In physical applications, the mapping on the side boundary may well be fixed in time.  It is also possible to consider more general $B_z$ distributions on the boundaries $z=0,1$, providing ${\\bf A}$ satisfies \\eqref{eqn:gcon} for an appropriate ${\\bf A}^{\\rm ref}$. In that case, the canonical 1-form $\\alpha$ is $\\alpha=rA_\\phi(r,\\phi,z)\\,\\mathrm{d}\\phi$, and the possible mappings $G$ that preserve $\\alpha$ (and are undetected by ${\\cal A}$) may differ, although they take the general form of a cotangent lift $G(\\phi_0,p)=\\big(T(\\phi_0),p(\\mathrm{d}T/\\mathrm{d}\\phi_0)^{-1}\\big)$, where $p=rA_\\phi$ is the generalized momentum (see Ref. \\onlinecite{marsden1994}, Proposition 6.3.2). If $\\widetilde{F}=F$ on the side boundary $r=R$, it follows that $T(\\phi_0)=\\phi_0$ and hence again $G={\\rm id}$.  Finally, one could consider a toroidal domain where ${\\bf B}$ and ${\\bf v}$ are periodic in $z$ and remove the condition that ${\\bf v}= 0$ on the boundaries $z=0,1$. In that case, the freedom to shuffle around field line endpoints on $z=0,1$ means that topological equivalence is a weaker notion, although it is certainly not true that any two field line mappings are equivalent, as would be the case if one had the freedom to apply independent motions on both boundaries. In the periodic case, the cross-section $z=0,1$ is no longer distinguished by the boundary conditions. Changing cross-section has the same effect on $F$ as an ideal evolution, so that the new mapping may be written $\\widetilde{F}=S\\circ F\\circ S^{-1}$ for some field line mapping $S$. If ${\\cal A}$ is the flux function for $F$, then one can show that the flux function $\\widetilde{\\cal A}$ for $\\widetilde{F}$ is given by \\begin{equation} \\widetilde{\\cal A} = (S^{-1})^*({\\cal A} + F^*\\chi - \\chi), \\end{equation} where $\\chi$ is related to the mapping $S$ by $S^*\\alpha - \\alpha = \\mathrm{d}\\chi$ (see Ref. \\onlinecite{haro1998}). The practical problem of determining whether two given topological flux functions are related in this way remains for further investigation."
    },
    "1208/1208.1766_arXiv.txt": {
        "abstract": "We characterise the typical offset between the Dark Matter (DM) projected centre and the Brightest Cluster Galaxy (BCG) in 10,000 SDSS clusters. To place constraints on the centre of DM, we use an automated strong-lensing (SL) analysis, mass-modelling technique which is based on the well-tested assumption that light traces mass. The cluster galaxies are modelled with a steep power-law, and the DM component is obtained by smoothing the galaxy distribution fitting a low-order 2D polynomial (via spline interpolation), while probing a whole range of polynomial degrees and galaxy power laws. We find that the offsets between the BCG and the peak of the smoothed light map representing the DM, $\\Delta$, are distributed equally around zero with no preferred direction, and are well described by a log-normal distribution with $\\langle \\log_{10}(\\Delta~[h^{-1} $Mpc$])\\rangle=-1.895^{+0.003}_{-0.004}$, and $\\sigma=0.501\\pm0.004$ ($95\\%$ confidence levels), or $\\langle \\log_{10}(\\Delta~[\\arcsec])\\rangle=0.564\\pm0.005$, and $\\sigma=0.475\\pm0.007$. Some of the offsets originate in prior misidentifications of the BCG or other bright cluster members by the cluster finding algorithm, whose level we make an additional effort to assess, finding that $\\sim10\\%$ of the clusters in the probed catalogue are likely to be misidentified, contributing to higher-end offsets in general agreement with previous studies. Our results constitute the first statistically-significant high-resolution distributions of DM-to-BCG offsets obtained in an observational analysis, and importantly show that there exists such a typical non-zero offset, in the probed catalogue. The offsets show a weak positive correlation with redshift, so that higher separations are generally found for higher-$z$ clusters in agreement with the hierarchical growth of structure, which in turn could potentially help characterise the merger, relaxation and evolution history of clusters, in future studies. In addition, the effective DM centre we adopt here, namely the peak of the smoothed light distribution representing the DM, can constitute a natural and alternative definition of cluster centers for optically-selected cluster catalogues. ",
        "introduction": "Galaxy clusters are the largest gravitationally-bound objects in the Universe, and are formed at later stages and relatively lower redshifts, in the hierarchical model. As such, galaxy clusters can shed light on the high end of the cosmic mass function and the evolution history of the Universe, and probe the acceptable cosmological model \\citep[e.g.][]{Allen2011clusterparareview}.  Digitised large sky surveys such as the \\emph{Sloan Digital Sky Survey} (SDSS; see \\citealt{York2000SDSS_tech,Abazajian2003SDSS_I,Abazajian2009SDSS_7}), and increasing computational power, have driven in recent years statistical analyses of extensively large cluster samples. In these (mainly optical imaging) data, galaxy clusters are identified usually in an automated manner, via dedicated finding algorithms (e.g. \\citealt{Postman1996Cat+finder,Kepner1999Finder,Gal2000Finder,GladdersYee2000Finder,Koester2007maxBCG_finder} and references therein; \\citealt{Pierpaoli2011SDSScatalog}). In addition, recent developments allow an independent detection of unprecedentedly large numbers of galaxy clusters in the X-ray and Sunyaev-Zeldovich Effect (SZE) observations \\citep[e.g.][]{Bohringer2004ClusterCat,Ebeling1998ClusterCat,SPT2010ClustersV1,SPT2011Clusters_V2,Planck2011ClustersV1}. Very large cluster samples have been used in many recent studies such as analyses of weak and strong lensing \\citep[e.g.][]{Johnston2007WL_SDSS,Mandelbaum2008StackedWL,Zitrin2011d}, or to place constraints on the cosmological parameters \\citep[e.g.][]{Rozo2010CosmoConstraintsSDSS}, and establish various scaling relations (e.g. \\citealt{Rozo2009MassRichSDSS,Planck2011XSZ_scaling,Bauer2012massrichness}; see also \\citealt{KravtsovBorgani2012Review}) as a few examples.  One of the major factors of noise or systematic uncertainty in these studies, especially when the clusters are identified optically and when the analysis is dependent on a predetermined centre, such as a stacked lensing analysis, cluster-background galaxy cross-correlation measurement (\\citealt{Johnston2007offset,Mandelbaum2010WL}; see also \\citealt{Oguri2010_25clusters,Umetsu2011b}), is the ``misentring'' of the BCG with respect to the dark matter (DM). Such offsets may result from either a misidentification of the cluster finding algorithm, or by a real measured offset between the projected DM centre and the BCG \\citep[see also][]{Johnston2007offset}, and will have a smoothing effect on the, e.g., stacked lensing signal (see \\citealt{OguriTakada2011}). A unique way to trace the offsets of the correctly identified clusters (hereafter, for simplicity, we dub these \\emph{true}, or \\emph{real} offsets to distinguish them from the misidentified clusters), to high accuracy, is by a strong-lensing (SL) analysis, where usually the multiple images in the core are used to accurately map the galaxies and dark mass distributions \\citep[e.g][]{Zitrin2009b,Umetsu2011}.  Due to the importance of the miscentring effect, there have been recent efforts to characterise its amplitude and size distribution. For example, \\citet{Johnston2007offset} quantified this effect by using N-body simulation-based mock galaxy catalogues, and then running their cluster-finding algorithm (maxBCG in that case; \\citealt{Koester2007maxBCG_finder}) to identify clusters in these catalogues, comparing the resulting BCG positions to the centres of the DM halos in the input simulations. They found that the offsets, i.e., for the clusters that are not centred on their BCG, are well described by a 2D gaussian of the form $P(R_{s})=\\frac{R_{s}}{\\sigma^{2}} \\exp(-\\frac{1}{2} (\\frac{R_{s}}{\\sigma})^{2})$, with $R_{s}$ being the magnitude of the offset and $\\sigma=0.42~ h^{-1}$ Mpc, and that the effect they traced is indeed dominated by misidentified BCGs. More recently \\citet{HilbertWhite2010} similarly found $\\sigma=0.34-0.41~ h^{-1}$ Mpc (for different WMAP cosmologies). \\citet{Johnston2007offset} also found that the fraction of misidentified clusters decreases with richness: $\\sim60\\%$ of the poorer ($N_{gal}\\sim10$) clusters are correctly identified, versus $\\sim85\\%$ of the richer ($N_{gal}\\sim100$) clusters. Although using these numerical simulations one can estimate the level and distribution of misidentified clusters, in such procedures the galaxies are assigned to DM halos in the simulation to begin with, so that the \\emph{true} offset distribution, i.e. for clusters whose BCG was correctly identified, is hard to assess.  Observationally, several studies also examined the offsets between the BCG and the X-ray peak (or centroid), or recently, also with the SZ peak. For example, \\citet{LinMohr2004} listed the BCG to X-ray peak offsets for a few dozen Two Micron All-Sky Survey (2MASS, \\citealt{Jarrett2000_2mass}) clusters. They found, that about $75\\%$ of the identified clusters lay within $0.06r_{200}$, and $90\\%$ within $0.38r_{200}$, with a $\\sim10\\%$ contamination level of possibly misidentified BCGs. They suggested, that given these high fractions of small X-ray to BCG separations, the timescale for the BCG to sink to the cluster potential minimum may be short compared to the relaxation timescale of the intracluster gas. Or alternatively, the scale of merger required to perturb the X-ray properties of the cluster could be smaller than the scale of the merger required to offset the BCG from the cluster centre \\citep{LinMohr2004}. \\citet{MannEbeling2012} examined the X-ray peak and centroid offsets from the BCG, in 108 of the most luminous X-ray clusters, with the goal of constraining the evolution with redshift of the cluster merger fraction, so that they also characterised the evolution of such offsets in redshift. They found that the distribution is roughly log-normal, and centred at 11.5 and 21.2 kpc for the offset of the BCG from the X-ray peak, and from the X-ray centroid, respectively. In addition, they found an evolution for these offsets with redshift, so that higher separations are generally expected for higher redshift clusters, probably as a sign of higher merging fraction. \\citet{Sehgal2012offsetsSZBCG} recently examined the SZ signal for 474 optically-selected (maxBCG) clusters and 52 X-ray selected (MCXC) clusters using data from the Atacama Cosmology Telescope (ACT). For the optically-selected sample, they found that the \\emph{Planck} and ACT measurements could be explained if one assumed that the BCGs are offset from the SZ peaks uniformly between 0 and 1.5 Mpc. However, they point out that other factors (rather than the BCG offset) could be in play in explaining the observed discrepancies, especially since for the X-ray-selected sample, a much narrower BCG to X-ray peak offset distribution was found, peaking within 0.2 Mpc. In recent work, also, \\citet{Song2012OffsetsSZ} find that the BCG to SZ centroid offset distribution, in 146 South Pole Telescope selected clusters with BCGs well identified in follow-up optical and near-infrared observations, is similar to that found previously in X-ray samples.  \\citet{Oguri2010_25clusters} observationally examined the offset of the BCG from the centre of mass obtained in weak lensing (WL) analyses of a sample of 25 clusters. They found that the DM centre is overall consistent with that of the BCG (within $2\\sigma$ level), and that the observed distribution can be described by two components. The first, significant component describing the small offsets, is a 2D Gaussian with $\\sigma=0.09~ h^{-1}$ Mpc, and the second less-significant component describing the tail of larger separations, is fitted by the \\citet{Johnston2007offset} finding: a 2D Gaussian with $\\sigma=0.42~ h^{-1}$ Mpc. They concluded that about 10\\% of their sample are BCG-offsetted clusters. \\citet{Oguri2010_25clusters} also probed the same effect by complementary X-ray fits for their sample, but found no specific correlation as the X-ray centroids are usually more closely centred on the BCG (e.g. \\citealt{LinMohr2004,Maughan2008clusterEv,MannEbeling2012}, although see also \\citealt{Shan2010_38offsets}). They also note that the typical error on the mass centroid measurement in their WL analysis is $\\sim50~ h^{-1}$ kpc, and there are a few clusters in their sample that have errors larger than $100~ h^{-1}$ kpc, which are non-negligible compared to the widths of the resulting 2D Gaussian distributions. In addition, \\citet{Dietrich2012WLoffsets} recently found, by thoroughly simulating WL observations, that generally, the magnitude of peak-offsets in WL maps could primarily be a direct result of shape-measurement noise and smoothing of the WL mass maps. It is thus clear that for a more quantitative assessment of the distribution of BCG offsets, a much larger sample is needed, probed with sufficiently high resolution.  Recently, \\citet{George2012offsets} examined in 129 galaxy groups (halo masses of up to $10^{14}~M_{\\odot}$) in the COSMOS field, the offset between the WL halo centre and other tracers such as the BCG, most massive group galaxy, or X-ray centroid. They found that the BCG is one of the best tracers of the centre-of-mass, and is offset typically by less than 75 kpc from the (dark) halo centre. In addition, they concluded that similar offset distributions are highly susceptible to the centre definitions, i.e., centres defined following certain intensity centroids, can largely differ from those defined via the corresponding intensity peaks \\citep[see][]{George2012offsets}, and that not accounting for the miscentring effect can cause a 5-30\\% bias in stacked WL analyses. The X-ray to BCG offset magnitude is generally consistent with several previous studies mentioned therein \\citep[e.g.][finding typically a few dozen kpc offsets or less]{Sheldon2001offsets,Koester2007maxBCG_cat,Sanderson2009locuss}.  \\citet{Shan2010_38offsets} characterised the offsets between the X-ray peaks and lensing centres in 38 clusters (see also \\citealt{Allen1998L-XrayDisc}). Although most clusters show small offsets, as is also usually seen in such lensing-analyses centre to BCG comparison \\citep[e.g.][]{Smith2005,Richard2010locuss20}, about $45\\%$ of their clusters, usually the merging, multiple-clump ones, show larger separations than $10\\arcsec$, with a maximum of $\\simeq54\\arcsec$ (or $\\sim200$ kpc, see also \\citealt{Forero-Romero2010}). This, however, may be a result of either large fractions of unrelaxed clusters in their sample ($\\sim60\\%$), and, the ensemble of different SL techniques used for the comparison, many of which pre-assume or iterate for the DM centre while adopting a symmetric DM distribution (see references therein), which may be unrealistic given the perturbed and complex matter distribution seen especially in unrelaxed clusters. In that sense, such offsets or even the known discrepancy between mass estimates from lensing and X-ray \\citep[e.g.][]{Allen1998L-XrayDisc,Richard2010locuss20}, may not be surprising \\citep[see][]{Shan2010OriginOffsets}.  \\begin{figure} \\centering \\includegraphics[width=90mm]{NvsS.eps} \\caption{Effect of the prior smoothing polynomial degree $S$ on the posterior distribution of offsets between the BCG assigned by the GMBCG catalogue and the DM centre (for fixed $q=1.2$). As higher $S$ values entail usually steeper mass profiles, and have more degrees of freedom to describe the substructure, higher $S$ values entail generally smaller BCG-DM offsets, converging towards the BCG near the boundary of our chosen range, $S=24$.} \\label{NvsS} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=90mm]{NvsQ.eps} \\caption{Effect of the prior galaxy surface-density power-law $q$ on the posterior distribution of offsets (for fixed $S=12$). Compared to the effect of the polynomial degree seen in Figure \\ref{NvsS}, the effect of the galaxy power law $q$ on the resulting distribution and the location of the peak, is negligible.} \\label{NvsQ} \\end{figure}   In a recent work, \\citet{Einasto2012} studied substructure and multimodality in groups and clusters of galaxies in SDSS DR8, using several tests to characterise the member distribution. They found that the distribution of distances from the cluster (or component) centre for the corresponding brightest galaxies shows that most of these galaxies are located preferentially close to the (sub)cluster centre (typically of an order of $\\sim0.1$ $h^{-1}$ Mpc scales, see Fig. 4 therein), although substantial numbers can show also higher separations (typically of a $\\sim1$ $h^{-1}$ Mpc scale).  Here, we aim to characterise the typical offset between the DM peak and the BCG of a statistically-significant sample, with a successful mass-modelling tool for SL analyses \\citep{Zitrin2009b} which does not require the DM centre to be predetermined, nor assumes a symmetric or any particular pre-known shape for the DM distribution. This method was recently adapted for automated use on 10,000 SDSS clusters, deducing the first observational, universal distribution of Einstein radii \\citep{Zitrin2011d}. The advantage of this approach in evaluating the BCG offsets is that the resolution is very high (an order of 0.1\\arcsec), and that the analysis can be performed blindly, on a very large sample. The results of this study will constitute the first statistically-significant, observational measure of this effect, and supply a complementary measure to compare to the previous studies mentioned above. In addition, the offset between the DM peak and the BCG may eventually help characterise the evolution of cluster galaxies and their host clusters, as more relaxed clusters can be anticipated to show decreasing or negligible offsets.  We analyse here the same 10,000 clusters from \\citet{Zitrin2011d}, which were randomly drawn from the Gaussian Mixture Brightest Cluster Galaxy (GMBCG; \\citealt{Hao2010GMBCG_cat}) SDSS cluster catalogue. In practice, these span the full redshift and richness ranges covered by the full catalog. The clusters had been found using the Error Corrected Gaussian Mixture Model algorithm \\citep{Hao2009GMBCG_find} to identify the BCG plus red sequence feature, convolving the identified red sequence galaxies with a spatial smoothing kernel to measure the clustering strength of galaxies (within 0.5 Mpc) around BCGs. The technique was applied to the Data Release 7 of the Sloan Digital Sky Survey and produced a catalogue of over 55,000 rich galaxy clusters in the redshift range $0.1 < z < 0.55$. The catalogue is approximately volume limited up to redshift $z\\sim0.4$ and shows high purity and completeness when tested against a mock catalogue, and when compared to other well-established SDSS cluster catalogues such as maxBCG (\\citealt{Koester2007maxBCG_cat}; for more details see \\citealt{Hao2010GMBCG_cat}).  \\begin{figure} \\centering \\includegraphics[width=80mm]{2DdistRealpixscalev3.eps} \\caption{2D (posterior) distribution of offsets between the BCG assigned by the GMBCG catalogue and the DM centre. As can be seen and expected, the offsets are distributed equally around zero with no preferred direction.} \\label{2dDist} \\end{figure}  The paper is arranged as follows: in \\S 2 we detail the method incorporated in order to obtain the BCG-DM offsets. In \\S 3 we report the results and factors of uncertainty, which are further discussed in \\S 4 and concluded in \\S 5. Throughout we use a standard $\\Lambda$CDM cosmology with ($\\Omega_{\\rm m0}=0.3$, $\\Omega_{\\Lambda 0}=0.7$, $h=0.7$), and distances are usually given in $h^{-1}$ Mpc, to ease the comparison to other works. To avoid possible confusion, we also note that all logarithmic quantities and syntax in this work are in base 10, unless stated otherwise, and are denoted equivalently as either ``$\\log_{10}$'' or ``\\emph{Log}''. ",
        "conclusions": "\\label{discussion}  Although only relatively little work was conducted on the subject, the DM to BCG offsets have been increasingly examined in recent years, due to growing technical and observational capabilities enabling statistical studies of large samples, such as stacked WL analyses, and since a knowledge of the DM-BCG typical offset plays an important role in them. Here we further discuss the results of our work and other possible effects of uncertainty.  In \\S \\ref{results}, and Figures \\ref{1dDistLog} and \\ref{1dDistLogarcsec}, we presented the measured distribution of BCG-DM offsets. An important question that arises given the measured distribution, is whether the BCGs are on average centred in their cluster potentials (described by the smoothed light distribution) with some finite deviation, or whether there is any evidence for a significant offset. To address this question, we perform a student-t test, which tests the hypothesis that a set of random variables is drawn from a Gaussian with either known or unknown mean and variance. The specific hypothesis tested in our case is whether the measured offsets can be considered as drawn from a Gaussian with zero mean and arbitrary variance. Suppose the projected BCG coordinates relative to their cluster centres are indeed random variates drawn from a two-dimensional Gaussian distribution, then their radial distances to the cluster centres, which we have measured here, are clearly not distributed in a Gaussian way. However, we can then return the radial distances into a set of Gaussian random variates by multiplying with the sine or cosine of a random phase with a flat distribution between $[0, 2\\pi]$. This multiplication with a random phase removes the geometrical effect that the differential area shrinks proportional to the distance from the centre, which makes radial distributions peak at finite radii. This corresponds to projecting the radial distance on an arbitrarily oriented coordinate axis. Thus, if our hypothesis is true, the radial distances multiplied with the sine or cosine of a random phase angle are Gaussian variates $\\{x_i\\}$ with zero mean. To this set of numbers, the student-t test can now be directly applied.  Its normalised test statistic, \\begin{equation} T(x) = \\frac{\\sqrt{N}\\bar x}{s(x)}\\;, \\end{equation} where $\\bar x$ and $s(x)$ are the mean and the standard deviation of the set $\\{x_i\\}$, follows a student-t distribution with $N-1$ degrees of freedom if the hypothesis is true. For our sample, we clearly have to reject this hypothesis. While the set $\\{x_i\\}$ can be well described by a Gaussian distribution, this distribution does not have zero mean. We obtain that the absolute t-statistics is 0.1(0.19), with an error probability of 0.42(0.46), that the linear(logarithmic) distribution is drawn from a Gaussian with a zero mean. We therefore conclude that although relatively small, the typical offset is significantly, non-zero.   In Figures \\ref{histvsz} and \\ref{histvszarcsec} we plotted the offset distribution in different redshift bins. The evolution of the mean and width of the (log-normal) distributions as a function of bin redshift is summarised in Figures \\ref{growthMPC} and \\ref{growthArcsec}. As can be seen therein, the mean of both the (log-normal) distribution in physical scales and in angular scales, increases steadily with redshift. The observed evolution, is however not significant: it is of the order of $1\\sigma$ for physical scale distribution, and only half a $\\sigma$ for the angular scale offset distribution. This low significance renders these trends at best, tentative. We note, however, that in a recent work \\citep{Zitrin2011d}, in which we analysed the Einstein radius distribution of the same GMBCG catalogue, we also uncovered similar (insignificant but monotonic) evolvement in redshift: throughout the same (volume-limited) redshift range, the mean Einstein radius decreases continuously with redshift. If real, these evolvements in redshift, in both the Einstein radius and the DM-BCG offset distributions, could together help characterise the evolution, relaxation, and merger history of galaxy clusters more generally, in addition to other complementary studies.  As mentioned (\\S \\ref{intro}), \\citet{Oguri2010_25clusters} observationally examined the offset of the BCG from the centre of mass obtained in weak lensing (WL) analyses of a sample of 25 clusters. They found that the DM centre is overall consistent with that of the BCG (within $2\\sigma$ level), and that the observed distribution can be described by two components. The first, significant component describing the small offsets, is a 2D Gaussian with $\\sigma=0.09~ h^{-1}$ Mpc, and the second less-significant component describing the tail of larger separations, is fitted by the \\citet{Johnston2007offset} finding \\citep[see also][]{HilbertWhite2010}: a 2D Gaussian with $\\sigma=0.42~ h^{-1}$ Mpc. Note that in our work we do not characterise the offset distribution of the misidentified clusters, but only the correctly identified ones. In addition, we cannot explicitly compare to the results of \\citet{Oguri2010_25clusters}, since in their work the centres are those determined by fitting a symmetric DM distribution (e.g., NFW) to the overall, larger-scale WL data (this in fact was recently found to be problematic especially in merging clusters, see \\citealt{George2012offsets}), while in our work we simply measure the range of possible locations of the central DM peak.  \\citet{Shan2010_38offsets} characterised the offsets between the X-ray peaks and lensing centres in 38 clusters (see also \\citealt{Allen1998L-XrayDisc}). Although most clusters show small offsets, as is also usually seen in such lensing-centre to BCG comparison \\citep[e.g.][]{Smith2005,Richard2010locuss20} showing typically less than $5\\arcsec$ offsets compatible with our results here, about $45\\%$ of their clusters and especially the merging, multiple-clump ones, show larger separations than $10\\arcsec$, with a maximum of $\\simeq54\\arcsec$. This, however, may be a result of either large fractions of unrelaxed clusters in their sample ($\\sim60\\%$), and, the ensemble of different SL techniques used for the comparison, many of which pre-assume or iterate for the DM centre while adopting a symmetric DM distribution (see references therein), which may be unrealistic given the perturbed and complex matter distribution seen especially in unrelaxed clusters. In that sense, such offsets or even the known discrepancy between mass estimates from lensing and X-ray \\citep[e.g.][]{Allen1998L-XrayDisc,Richard2010locuss20}, may not be surprising \\citep[see][]{Shan2010OriginOffsets}. This in fact is a crucial point to make here, as we implied above: In our work we do not fit to the data a symmetric DM distribution for which the effective DM centre may be in practice different than the DM peak, whose offset from the BCG we characterise here.  Several studies have also examined the offsets between the BCG and the X-ray peak or centroid \\citep[e.g.][]{LinMohr2004,Maughan2008clusterEv}. Recently, \\citet{MannEbeling2012} examined the X-ray peak and centroid offsets from the BCG, in 108 of the most luminous X-ray clusters, with the goal of constraining the evolution with redshift of the cluster merger fraction, so that they also characterised the evolution of such offsets, with redshift. Similar to our result, they also found that the distribution is (roughly) log-normal, and centred at 11.5 kpc ($H_{0}=70$ km~s$^{-1}$~Mpc$^{-1}$) for the offset of the BCG from the X-ray peak, overall similar to, or of the same order of, the peak centre we find for our BCG-DM offset distribution: 12.7 $h^{-1}$ kpc. For the BCG offset from the X-ray centroid, their peak is $\\sim$twice as large. In addition, they found an evolution for these offsets with redshift, so that higher separations are generally expected for higher redshift clusters, probably as a sign of higher merging fraction, similar to the (tentaive) evolution we observe here between the DM peak and the BCG. This evolution generally agrees also with other complementary studies such as (a different X-ray sample) brightness centroid shifts, metallicity, brightness profile steepness, or other similar relations found in cluster evolution works \\citep[e.g.][and references therein]{Maughan2008clusterEv}.  A similarly interesting question which should be investigated quantitatively, is the correlation between the offset magnitude and the degree of relaxation. While in our work the BCG to DM peak offset may imply, generally, the degree of relaxation, other independent relaxation measures would be needed in order to obtain at least a reference sample for one to find a correlation between these. In a similar manner, the degree of relaxation have been defined by various (simulation-based) works in the literature \\citep[e.g.][]{Thomas200tauCDM,Neto2007}, where the offset between the potential minimum and the centre of mass, constitutes one of the measures for relaxation \\citep[see][]{Lemze2011AnisProfiles}, however these require knowledge of both. It would be therefore worth developing other independent relaxation measures to compare with, in future studies.  In a similar context, and although in our previous analyses we have successfully covered large ranges of these parameters, once many more clusters are analysed in detail using multiple images, in combination also with WL data where possible, it would be interesting to test whether the procedure described here still applies to new limits of richness, luminosity, relaxation, concentration, mass, et cetera.  \\begin{figure} \\centering \\includegraphics[width=90mm]{GrowthOfMuWithZ.eps} \\caption{Growth of the log-normal distribution mean, $\\langle \\log_{10}(\\Delta [h^{-1} $Mpc$])\\rangle$ (\\emph{open circles}), and width, $\\sigma$ (\\emph{vertical error bars}), as a function of redshift. The \\emph{horizontal error bars} represent the redshift bin width. Although it is insignificant, i.e. of the order of $1\\sigma$, a tentative trend is seen as a function of redshift, so that higher separations and lower widths are seen for higher-$z$ clusters. The linear and quadratic least-squares fit to the data are given in the upper left corner, and overplotted as dash-dotted red and black lines, respectively.} \\label{growthMPC} \\end{figure} \\begin{figure} \\centering \\includegraphics[width=90mm]{GrowthOfMuWithZArcsec.eps} \\caption{Same as Figure \\ref{growthMPC}, but for the log-normal distribution in arcsecond binning. This Figure, along with Figure \\ref{growthMPC}, shows that there exists a tentative (of the order of only half a $\\sigma$ here) but steady trend possibly originating from the evolution and merger history of clusters, so that higher redshift and therefore less relaxed clusters tend to show (both physical-distance and angular) higher separations, and lower (log-normal) distribution widths. One way to interpret the minor distribution-width trend, is that higher redshift clusters exhibit a somewhat more coherent population, say, in terms of relaxation, than lower redshift clusters that show a wider variety of relaxed and unrelaxed clusters.} \\label{growthArcsec} \\end{figure}  We also note, that an alternative smoothing process to consider, could be based on a Gaussian filter smoothing rather a fit to a 2D polynomial as performed here, since the latter could suffer from various edge effects and other boundary condition artefacts (although, importantly, with no noticeable effect on the central, DM peak location; the length scale is an order of magnitude different). The reason we did not use a Gaussian smoothing, but the 2D spline interpolation, is that the latter method is well vetted and works remarkably well in reconstructing many mass distributions, as seen in our previous lensing analyses mentioned throughout. We are, however, in the process of examining the Gaussian smoothing alternative, with the goal of establishing the range of relevant Gaussian-filter widths to be implemented (e.g. as priors) in the future analyses, for comparison.  In our work here we made an additional attempt to identify the level of misidentified and grossly miscentred clusters. We found that substantial part of the misidentified clusters show a typical offset behaviour in the $q$-$S$ parameter plane, of which we made use to exclude such clusters from the fit and estimate their portion in the whole sample. Accordingly, we found that about $10\\%$ of the clusters in the catalogue may be misidentified, which is \\emph{smaller} but of the same general order as found in the previous studies mentioned throughout. In that sense, also, note that we do not attempt to characterise the offset distribution of these misidentified clusters, nor can we distinguish them from real, grossly miscenetred clusters, if these indeed exist and survived our cleaning procedure. Both by the offset log-normal shape we find here which drops for higher offsets, and the difference between the $10\\%$ noise level we find and the level found in other studies, the effect of both largely offsetted and misidentified clusters in the remaining sample (after cleaning the misidentified clusters we spotted in our procedure) on our results, is expected to be minimal (see also \\S \\ref{offsetNoiseS1})."
    },
    "1208/1208.1285_arXiv.txt": {
        "abstract": "We present a detailed analysis of the spatially and spectrally resolved $^{12}$CO $J$=2$-$1 and $J$=3$-$2 emission lines from the TW Hya circumstellar disk, based on science verification data from the Atacama Large Millimeter/Submillimeter Array (ALMA).  These lines exhibit substantial emission in their high-velocity wings (with projected velocities out to 2.1\\,km s$^{-1}$, corresponding to intrinsic orbital velocities $>$20\\,km s$^{-1}$) that trace molecular gas as close as 2\\,AU from the central star. However, we are not able to reproduce the intensity of these wings and the general spatio-kinematic pattern of the lines with simple models for the disk structure and kinematics.  Using three-dimensional non-local thermodynamic equilibrium molecular excitation and radiative transfer calculations, we construct some alternative models that successfully account for these features by modifying either (1) the temperature structure of the inner disk (inside the dust-depleted disk cavity; $r < 4$\\,AU); (2) the intrinsic (Keplerian) disk velocity field; or (3) the distribution of disk inclination angles (a warp). The latter approach is particularly compelling because a representative warped disk model qualitatively reproduces the observed azimuthal modulation of optical light scattered off the disk surface.  In any model scenario, the ALMA data clearly require a substantial molecular gas reservoir located inside the region where dust optical depths are known to be substantially diminished in the TW Hya disk, in agreement with previous studies based on infrared spectroscopy.  The results from these updated model prescriptions are discussed in terms of their potential physical origins, which might include dynamical perturbations from a low-mass companion with an orbital separation of a few AU. ",
        "introduction": "In the standard model for star formation, the collapse of a slowly rotating molecular cloud core at least partially conserves its angular momentum by forming a compact, flattened disk that channels mass onto a central protostar \\citep{cassen81,terebey84}.  Once the remnant core material is accreted or dispersed, the gravitational potential of the star determines the kinematic properties of its disk.  At that time, the intrinsic disk velocity field should be described well by a Keplerian pattern of differential rotation in circular orbits, $v_{\\phi}(r) = v_k = \\sqrt{G M/r}$, where $\\phi$ is the azimuthal direction in a polar coordinate system, $G$ is the gravitational constant, $r$ is the distance from the star, and $M$ is the mass enclosed within $r$.  In most cases, the disk-to-star mass ratio is low enough that $M \\approx M_{\\ast}$ is a reasonable approximation.  Imaging spectroscopy data taken with millimeter interferometers first verified these signatures of disk rotation using the optically thick emission lines of the carbon monoxide (CO) molecule \\citep{koerner93,dutrey94,koerner95,mannings97}.  Subsequent work demonstrated that the rotation patterns were indeed Keplerian ($v_{\\phi} \\propto r^{-0.5}$) and could be used to make dynamical estimates of central star masses \\citep{guilloteau98,dutrey98,simon00}.  In practice, spectral imaging observations measure the intrinsic velocity field projected onto the sky, $v_{\\rm obs} = v_{\\phi} \\sin{i}$, where $i$ is the angle between the disk rotation axis and the observed line of sight (e.g., $i = 0\\degr$ corresponds to face-on).  A Keplerian velocity field can be exploited to probe gas at small radii in the disk, even if the angular resolution of the observations is relatively limited \\citep[e.g.,][]{dutrey08}.  This kind of physically motivated super-resolution can be achieved with a focus on the intensities (and ideally morphologies) in the high-velocity wings of emission lines.  For optically thick lines like those produced by the rotational transitions of CO, the line strength depends on the product of the gas temperature and emitting area \\citep[e.g.,][]{beckwith93}.  Since the projected emitting areas are typically small in the inner disk, emission in the line wings is generally weak; often well below the sensitivity thresholds for most millimeter-wave interferometers.  To further complicate issues, there are theoretical mechanisms and disk properties that could modify this simple prescription for interpreting the projected disk velocity field.  For example, in a very massive disk, gravitational instabilities could excite spiral waves that drive significant (radial) streaming motions.  This was suggested as a potential explanation for the complex kinematic environment of the AB Aur disk \\citep{lin06}, although others have associated its (possibly) non-Keplerian motions with envelope contamination \\citep{pietu05,corder05}.  Gas pressure gradients also lead to deviations from Keplerian orbits: the sub-Keplerian shift expected from thermal (hydrostatic) pressure is too small to observe directly \\citep{weidenschilling77}, but magnetic pressure might have a more substantial influence over gas motions \\citep[e.g.,][]{shu07,shu08}.  Alternatively, a disk might {\\it appear} to have non-Keplerian motions (even if $v_{\\phi} = v_k$ exactly) if its line-of-sight orientation ($i$) varies with radius -- that is, if the disk structure is warped.  Given that emission line wings are typically expected to be weak, disentangling these subtle (hypothetical) modifications to a simple projected velocity field from the intrinsic properties of the inner disk makes for a difficult task in data analysis.  Nevertheless, the potential for this unique access to the spatio-kinematic properties and physical conditions of the gas in the innermost regions of protoplanetary disks is extraordinarily compelling.  Perhaps the best opportunity to explore these features is with the massive, gas-rich disk around the nearest \\citep[$d \\approx 54$\\,pc;][]{vanleeuwen07} classical T Tauri star, TW Hya.  Aside from the enhanced sensitivity and spatial resolution afforded by its proximity, the TW Hya disk coincidentally has the added benefits of a nearly face-on viewing geometry \\citep[$i \\approx 6$-7\\degr;][]{krist00,qi04,hughes11}.  This means that the line wing emission generated by gas in the inner regions of the TW Hya disk has relatively small observed Doppler shifts from the systemic velocity.  Because the low disk inclination angle minimizes radial projection effects, subtle departures from a simple (unwarped, Keplerian) model for the projected velocity field should be more easily recognizable observationally.  In this article, we take advantage of the dramatically improved sensitivity now available in the Science Verification (SV) data products from the Atacama Large Millimeter/Submillimeter Array (ALMA) to help characterize the spatio-kinematic morphologies of the spatially and spectrally resolved CO $J$=2$-$1 and $J$=3$-$2 emission lines from the TW Hya circumstellar disk.  A brief overview of the data and its calibration are presented in \\S 2.  The data is investigated in detail in \\S 3, using emission line radiative transfer calculations and considering various toy models for the disk properties.  The results are synthesized in the larger context of the structure and kinematics in the inner regions of the TW Hya disk in \\S 4, and summarized in \\S 5. ",
        "conclusions": "We have conducted a detailed analysis of the $^{12}$CO $J$=2$-$1 and $J$=3$-$2 line emission from the gas-rich disk around the nearby young star TW Hya, making extensive use of the science verification data from the Atacama Large Millimeter Array (ALMA) commissioning effort.  With the substantial improvement in line sensitivity afforded by ALMA (even during its construction), we identified a novel feature in these CO spectra: the presence of relatively faint line wing emission extending out to high projected velocities of at least $\\pm2.1$\\,km s$^{-1}$ from the (systemic) line center, roughly twice the maximum velocities previously inferred for the TW Hya disk from less sensitive data \\citep{qi04,qi06,hughes11,andrews12}.  Given the nearly face-on viewing geometry of the TW Hya disk, gas on Keplerian orbits that generates these line wings must be located within a few AU from the central star.  This kinematic super-resolution was employed in an exploration of the temperature and velocity structure of the inner disk, utilizing a non-local thermodynamic equilibrium molecular excitation and radiative transfer code and a standard modeling prescription for circumstellar disks.  We found that the typical assumptions of a simple, Keplerian disk with a smooth gas temperature profile were unable to simultaneously produce both the CO line wings and the observed line emission morphology near the systemic velocity.  Instead, we developed three alternative model possibilities that are able to qualitatively account for the ALMA data by making small (parametric) adjustments to the spatio-kinematic structure of the gas in the inner disk.  In one such alternative, we considered a parametric scaling of the gas temperatures inside the known dust-depleted cavity at the disk center \\citep[$r < 4$\\,AU;][]{calvet02,hughes07} in an effort to boost the CO line wing intensities without substantially disturbing the emission morphology near the line core.  With the outer disk structure based on a fiducial model, an increase of the gas temperatures in the dust cavity by a factor of $\\sim$3 successfully accounts for the ALMA CO observations.  While such a temperature scaling might seem simplistic and artificial, it is in fact consistent with continuum radiative transfer models for the TW Hya disk: the low optical depths inside the disk cavity mean that stellar radiation is effectively unattenuated, leading to substantial heating \\citep{calvet02,andrews12}.  If we assume that gas and dust temperatures are correlated (if not identical), then the factor $\\delta T \\approx 3$ derived here is quite reasonable.  For reference, the models developed by \\citet{andrews12} have dust temperatures that increase by a factor of $\\sim$4 over a narrow ($\\sim$1\\,AU-wide) region around the cavity radius (see their Figure 3).  It is interesting to consider the key implication of this model: the CO emission line wings detected by ALMA may offer an indirect (and not very practical) observational signature of a low-density dust cavity, but they also {\\it require} that this cavity still hosts a significant reservoir of molecular gas.  Of course, the presence of gas inside the TW Hya disk cavity has been inferred from other diagnostics: accretion indicators \\citep{muzerolle00}, as well as emission lines of simple atoms and molecules \\citep{najita10}, including CO \\citep[e.g.,][]{rettig04}.  Based on a rotational diagram constructed from the infrared spectrum of CO fundamental ($v$ = 1$-$0) rovibrational emission lines, \\citet{salyk07} used a simple slab model to estimate $T = 650$-1500\\,K and $N_{\\rm co} = 0.2$-$1.4 \\times 10^{19}$\\,cm$^{-2}$ for an emitting region that corresponds to $r \\approx 0.1$-1\\,AU.  Taking the mean temperature from the two CO lines, we find $T = 640$-1660\\,K and $N_{\\rm co} = 14$-$140\\times 10^{19}$\\,cm$^{-2}$ for our ``hot\" model in the same region -- remarkably similar, keeping in mind that column density estimates from optically thick lines are inherently ambiguous.  However, other analyses of the CO rovibrational spectrum suggest much lower temperatures \\citep[by a factor of $\\sim$2;][]{rettig04,pontoppidan08,salyk09}.  Moreover, our ``hot\" model assumes enhanced temperatures over a continuous emitting area out to 4\\,AU: if a smaller area were adopted, the temperature scaling factor $\\delta T$ would have to be proportionately increased to compensate for the observed line wing intensities (notably to uncomfortable levels that start to impinge on the thermal dissociation energy).  Given the uncertainties involved in analyzing both kinds of CO data, it is not clear whether a ``hot\" model derived from the rotational lines (as in \\S 3.3.1) is necessarily in conflict (i.e., too hot) with the rovibrational spectrum.  In any case, it is interesting to note that these two complementary methods could eventually be used to obtain independent constraints on the gas properties inside the TW Hya dust disk cavity.  In the near future, ALMA can provide sensitive high angular resolution images of these same CO transitions that will resolve the spatio-kinematic morphology of their line wings.  That information will be crucial for locating the CO emitting area, enabling more direct comparisons between the infrared and millimeter-wave diagnostics.  While the first model described above was devoted to modifications of the {\\it structure} in the inner disk, the remaining two alternatives focused on adjustments to the disk {\\it kinematics}.  Rather than increasing the intrinsic line intensities from a fixed emitting area, these other models produced the observed CO line wings by increasing the effective emitting area itself.  In one such scenario, we modified the intrinsic velocity field in the inner disk with a simple parametric deviation from Keplerian orbits.  We found in \\S 3.3.2 that a model which incorporated super-Keplerian rotation in the inner disk successfully reproduced the observed CO line emission.  Although there is considerable uncertainty involved, our ``non-Keplerian\" model rotated $\\sim$30\\%\\ faster than Keplerian speeds at $r = 10$\\,AU, and $\\sim$80\\%\\ faster at 1\\,AU ($f \\approx 1.3$ and 1.8, respectively).  If this deviation was restricted to much smaller disk radii only, the corresponding scaling factor $f$ would need to be substantially larger to account for the observed CO line wings.  Admittedly, of the three alternative models explored here, this scenario is the least physically-motivated.  While in principle a (reverse) pressure gradient could accelerate the disk material, most models that consider such gradients instead call for sub-Keplerian motions \\citep[e.g.,][]{weidenschilling77,shu07,shu08}.  However, it is possible that non-azimuthal gas velocities could mimic the parameterization used here.  For example, a molecular outflow with high mass-loss rates, directed primarily in the vertical direction (a $v_z$ component) and launched from the inner disk, might play a role.  The TW Hya disk is thought to drive a substantial photoevaporative wind \\citep[e.g.,][]{pascucci11}, but the molecular properties of that flow have not yet been explored in sufficient detail to easily compare with observations.  Another example might be a velocity field with radial streaming motion (a $v_r$ component), perhaps directed along spiral arms \\citep[e.g.,][]{quillen06}.  No such structural features have yet been identified for the TW Hya disk; but there are similar disks that do exhibit this kind of asymmetry \\citep{weinberger99,augereau99,grady01,clampin03,fukagawa04,fukagawa06,muto12}.  Though these possibilities seem exotic, they remain in consideration until realistic, testable models for more complex disk velocity fields are developed.  The final alternative model we developed to account for the ALMA CO data also relied on modifying the {\\it observed} disk velocity field, $v_{\\rm obs} = v_{\\phi} \\sin{i}$.  In this scenario, we maintained the standard assumption of Keplerian orbits ($v_{\\phi} = v_k$) and instead permitted the projection factor $\\sin{i}$ to vary with radius.  In \\S 3.3.3, we developed a simple parametric description of a disk warp that simultaneously accounts for the CO line wings and the nearly face-on emission morphology near the systemic velocity.  That simple model has an inclination of $\\sim$4\\degr\\ in the outer disk and increases only slightly to $\\sim$10\\degr\\ at $r = 1$\\,AU; the mean line-of-sight viewing angle averaged over the disk area is comparable to the $\\sim$4-7\\degr\\ inferred previously \\citep{krist00,weinberger02,qi04,hughes11}. Warps have been directly imaged in scattered light for the edge-on (effectively gas-free) debris disks around $\\beta$ Pic \\citep{kalas95,heap00} and AU Mic \\citep{liu04,krist05}, and indirectly inferred from quasi-periodic photometric variations in optical/infrared light curves in other edge-on cases, notably KH 15D \\citep{chiang04} and AA Tau \\citep{bouvier99}.  Warped viewing geometries have also been speculated for a few additional cases, including the disks around GM Aur \\citep[from CO spectral imaging;][]{dutrey98,hughes09} and HD 100546 \\citep[from scattered light images;][]{quillen06b}.  In the specific case of the TW Hya disk, \\citet{roberge05} have cited potential evidence for a warp based on an azimuthal asymmetry in their wideband {\\it Hubble Space Telescope} ({\\it HST})/STIS coronagraphic images of optical light reflected off the disk surface.  They identify a roughly sinusoidal variation of scattered light with the disk position angle in a region $\\sim$1.3-1.6\\arcsec\\ from the central star, with a peak near P.A. = 234\\degr\\ and a trough near 54\\degr\\ (measured E of N, and notably roughly $\\pm$90\\degr\\ from the disk major axis orientation).  \\citet{roberge05} speculated that this azimuthal modulation in the brightness profile might be the consequence of a varying illumination pattern in the outer disk, generated by the shadowing effect of a warped disk geometry at much smaller radii.  This is effectively the same type of model we have employed to explain the ALMA observations of the CO lines, so naturally we were motivated to test this possibility against the behavior of the \\citet{roberge05} scattered light data.  To do so, we implemented a simplified dust structure based on the ``warp\" model described in \\S 3.3.3, assuming that the dust and gas are co-located outside a radius of 4\\,AU: to adhere to the dust modeling in the literature, we assume there is effectively no dust for $r < 4$\\,AU \\citep[e.g.,][]{calvet02,hughes07}.  We adopt the relative mass ratio, grain composition, and size distribution used in the disk atmosphere models of \\citet{andrews12}.  The three-dimensional Monte Carlo radiative transfer code {\\tt RADMC-3D}\\footnote{\\url{http://www.ita.uni-heidelberg.de/$\\sim$dullemond/software/radmc-3d/}} (v0.30: C.~P.~Dullemond) was then used to generate a high-resolution synthetic image at a wavelength of 0.8\\,$\\mu$m, assuming isotropic scattering.  Following \\citet{roberge05}, we calculated a normalized intensity profile in 20\\degr\\ azimuthal bins in an annular ring from $r = 70$-88\\,AU.  \\begin{figure}[t] \\epsscale{0.85} \\plotone{f06.eps} \\figcaption{A comparison of the (normalized) azimuthal profile of optical light scattered off the TW Hya disk surface from both {\\it HST}/STIS coronagraph data \\citep[{\\it red circles};][]{roberge05} and a prediction based on our ``warp\" model ({\\it black curve}; note that this is {\\it not} a fit to the scattered light data).  Both profiles were derived by averaging the intensities in 20\\degr\\ bins from an annulus extending from $r = 70$-88\\,AU.  The sinusoidal modulation is a natural consequence of shadowing by a warped structure in the inner disk (higher inclination angles at small radii), which produces a varying illumination pattern of the disk surface at larger radii.  \\label{fig:az_scat}} \\end{figure}  The results are shown in Figure \\ref{fig:az_scat} and demonstrate a striking resemblance to the azimuthal modulation pattern inferred from the {\\it HST}/STIS data \\citep[][see their Figure 9]{roberge05}.  Taken at face value, this remarkable agreement might be considered compelling independent evidence for a very modest warp in the TW Hya disk structure.  Again, this scattered light model is {\\it not} a fit to the {\\it HST}/STIS observations of scattering off dust grains, but rather the predicted behavior based on one model that is able to describe ALMA observations of the CO gas structure.  The sinusoidal pattern of the azimuthal profile is a result of material in the inner disk -- where the warp is enhanced -- shadowing dust at larger radii.  The phase of that pattern is set by the assumed warp (rotation) axis, which was arbitrarily aligned with the observed major axis of the disk projected on the sky (151\\degr\\ E of N).  However, in practice the phase depends only relatively weakly on the warp axis orientation: shifting the latter by $\\pm$20\\degr\\ leads to only modest changes to the pattern phase (i.e., the data are not quite able to quantify the warp orientation).  The pattern amplitude depends more sensitively on the adopted model parameters, as it is set primarily by the height of the optically thick dust layer in the inner disk.  If the inner disk were substantially cooler, and therefore had lower scale heights, the level of shadowing would decrease and the amplitude of the azimuthal scattering pattern would be diminished.  Indeed, if we instead adopt the cooler, midplane dust temperatures found by \\citep{andrews12} to set the dust scale heights using Equation 4, we find a modest ($\\sim$20\\%) decrease in the peak amplitudes that are more consistent with the scattered light data.  Perhaps more important in this case is the truncation of the dust distribution for $r < 4$\\,AU.  If we were to assume no dust depletion in the disk cavity, the shadowing of the outer disk would be much more pronounced, with a normalized amplitude of $\\sim$1 in the azimuthal variation of scattered light.  In essence, there is a substantial degeneracy between the spatial distribution of dust in the inner disk (i.e., the cavity size and depletion factor, as well as the local scale heights) and the parameters that describe the warp.  Nevertheless, for a set of simple assumptions grounded in the ALMA CO data \\citep[and the dust cavity from the SED and resolved images;][]{calvet02,hughes07,andrews12}, we find a warp model is remarkably consistent with a key feature of the \\citet{roberge05} scattered light data.  Warps have been postulated to originate from a wide range of phenomena, including gravitational instabilities \\citep[e.g.,][]{sellwood10}, intense radiation fields \\citep{armitage97}, star-disk magnetic interactions \\citep[e.g.,][]{terquem00,flaherty10}, and dynamical perturbations from a stellar flyby \\citep{quillen06b,nixon10}.  However, none of these seem particularly plausible in this case, given the known properties of TW Hya and its disk.  In principle, a massive planet embedded in the inner disk with an orbit inclined out of the disk plane could also generate the modest warp inferred here \\citep{lubow01,marzari09}.  While this scenario has been extensively explored for the $\\beta$ Pic debris disk \\citep{mouillet97,dawson11}, the high gas densities in the TW Hya disk could lead to the substantial damping of a planet inclination -- and therefore the warp structure -- on a relatively short (viscous) timescale \\citep{papaloizou95}.  However, given the large uncertainties on the disk viscosities and densities, as well as the range of potential companion properties (e.g., masses, orbits), there is still a reasonable likelihood that a companion could maintain the modest orbital inclination required to sustain a small warp over the lifetime of the TW Hya disk \\citep[e.g.,][]{bitsch11}. Moreover, other effects related to planet-planet scattering in the system could produce repeated dynamical excitations of the gas disk \\citep[e.g.,][]{thommes03}.  We would be remiss not to point out that there is an alternative explanation for the \\citet{roberge05} scattered light asymmetry: namely, a spatial variation in the scattering phase function of the grains.  This latter possibility might better account for the observed kinks in the {\\it radial} scattered light profiles, which incidentally our toy model does not reproduce well (although perhaps could, given a considerable modeling effort beyond the scope of this article).  Other studies of the TW Hya scattered light disk from {\\it HST} make no clear mention of an azimuthal asymmetry \\citep[e.g.,][]{krist00,weinberger02}.  The original WFPC2 discovery paper by \\citet{krist00} does note that the ``right\" (we assume west) side of the disk appears brighter (consistent with the Roberge et al.~asymmetry), but they suggest that the discrepancy could just as easily be explained by an instrumental or calibration artifact.  We should also note that \\citet{pontoppidan08} used infrared spectro-astrometry measurements of the CO fundamental rovibrational lines to determine an inclination of only $\\sim$4-5\\degr\\ in the inner $\\sim$0.1-0.5\\,AU of the TW Hya disk.  The data presented here are not directly sensitive to such small radial scales, but the \\citet{pontoppidan08} claim of a warped {\\it outer} disk (where $i$ is an increasing function of $r$) is clearly not consistent with the ALMA data. These inferred variations in inclination all could be compatible if the disk tilt is not monotonic with radius, but instead has some maximum inclination at a few AU.  In any case, a simple warp model is a compelling potential explanation for both the ALMA CO spectral images and the {\\it HST}/STIS scattered light data, but it is not necessarily the correct (or unique) solution.  In principle, all three of the alternative models that successfully explain the ALMA CO data could ultimately be produced by dynamical interactions between the disk and a companion located a few AU from the central star.  Nevertheless, despite the known dust depletion at small radii or potential evidence for a warp in this disk, there is not yet any firm evidence for a close companion to TW Hya.  The most definitive high-resolution contrast limits for this system rule out companion masses $>$0.014\\,M$_{\\odot}$ at 2-4\\,AU separations \\citep{evans12}, suggesting that any perturber would need to be substellar or planetary in nature.  A previous claim for a ``hot Jupiter\" \\citep{setiawan08} has been refuted \\citep{huelamo08,rucinski08}, and, although there is a suggestive asymmetry in new mid-infrared data that is consistent with a faint point source near the edge of the disk cavity, its interpretation remains uncertain \\citep{arnold12}.  Along with deeper companion searches in the near future, we expect that some of the techniques presented here could be used to help provide new, empirical constraints on how dynamical interactions affect the kinematic and structural properties of the TW Hya gas disk and others like it."
    },
    "1208/1208.6236_arXiv.txt": {
        "abstract": "Recent {\\it XMM-Newton} studies of X-ray variability in the hard states of black hole X-ray binaries (BHXRBs) imply that the variability is generated in the `standard' optically-thick accretion disc that is responsible for the multicolour black-body emission. The variability originates in the disc as mass-accretion fluctuations and propagates through the disc to `light up' inner disc regions, eventually modulating the power-law emission that is produced relatively centrally.  Both the covariance spectra and time lags that cover the soft band strongly support this scenario.  Here we present a comparative spectral-timing study of {\\it XMM-Newton} data from the BHXRB SWIFT~J1753.5--0127 in a bright 2009 hard state with that from the significantly fainter 2006 hard state, to show for the first time the change in disc spectral-timing properties associated with a global increase in both the accretion rate and the relative contribution of the disc emission to the bolometric luminosity.  We show that, although there is strong evidence for intrinsic disc variability in the more luminous hard state, the disc variability amplitude is suppressed relative to that of the power-law emission, which contrasts with the behaviour at lower luminosities where the disc variability is slightly enhanced when compared with the power-law variations.  Furthermore, in the higher-luminosity data the disc variability below 0.6~keV becomes incoherent with the power-law and higher-energy disc emission at frequencies below 0.5~Hz, in contrast with the coherent variations seen in the 2006 data.  We explain these differences and the associated complex lags in the 2009 data in terms of the fluctuating disc model, where the increase in accretion rate seen in 2009 leads to more pronounced and extended disc emission.  If the variable signals are generated at small radii in the disc, the variability of disc emission can be naturally suppressed by the fraction of unmodulated disc emission arising from larger radii.  Furthermore the drop in coherence can be produced by disc accretion fluctuations arising at larger radii which are viscously damped and hence unable to propagate to the inner, power-law emitting region. ",
        "introduction": "\\label{sec:intro}  Variability on a broad range of time-scales, ranging from milliseconds to hours, is a well-known characteristic of black hole X-ray binaries (BHXRBs). The short-term variability (time-scales up to a few hundred seconds) is strongly dependent on the spectral state of the object, being particularly enhanced in the so-called hard state, with rms amplitudes up to $\\sim$ 40 per cent of the average flux \\citep{belloni2005,remillard2006,munozdarias2011}.  In the hard state, the X-ray emission is dominated by a hard power-law component (typically, $\\Gamma \\sim$ 1.4 -- 2.1) accompanied by weaker multi-colour black-body emission associated with the accretion disc, with $kT_{\\rm in} < 0.5$ keV \\citep{miller2006,reis2010}.  While the source of black-body photons is strongly suspected to be the accretion disc, there is some controversy as to the origin and physical location of the power-law component, which could be produced by Compton scattering in a hot corona \\citep{malzac2009}, in a hot inner flow \\citep{zdziarski1998} resulting from the evaporation of the inner radii of the accretion disc (thus truncating the disc, see e.g. \\citealt{narayan1994}), or at the base of the radio-emitting jet that is observed during this state \\citep{markoff2005}. Notwithstanding the actual physical source of hard photons (i.e. $E\\gtrsim 2$ keV), there is general agreement that the hard photon-emitting region must be relatively central, within tens of gravitational radii from the central object in the bright hard state.  Previously, most X-ray variability studies of the hard state were conducted using instruments with hard X-ray sensitivity, e.g. the Proportional Counter Array (PCA) on the {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}), which limited the study of the variability to the hard, power-law component.  Due to the anti-correlation of rms variability amplitude with spectral hardness as a source transitions to the disc-dominated soft state (e.g. \\citealt{belloni2005}), a picture arose where the disc itself was intrinsically stable and variability is generated in the power-law emitting region, usually envisioned as an unstable hot inner flow \\citep{churazov2001,done2007}.  However, recent work \\citep{wilkinson2009,uttley2011} has used the soft X-ray response and timing capability of the EPIC-pn instrument on board {\\it XMM-Newton} to carry out the first spectral-timing studies of the disc in the hard state. This work has shown strong evidence, in the form of variability (`covariance') spectra and time-lags, that the disc is intrinsically variable, and that disc variations in fact precede (and likely drive) the power-law variations, at least on time-scales of seconds and longer.  The role of the disc as the driver of variability can be explained by invoking the presence of mass-accretion rate fluctuations that originate in the disc at particular radii and propagate inwards. Such fluctuations vary at the local viscous time-scale corresponding to the radius where they are produced \\citep{lyubarskii1997,kotov2001,arevalo2006}. Because the $\\dot{m}$ variations at inner radii are modulated by the outer $\\dot{m}$ via a multiplicative process in the context of this model, variability is seen at all frequencies that correspond to the viscous time-scales of the radii generating the propagating signals. Once the perturbations reach the inner regions of the system, these modulate the power-law emission that accounts for most of the hard X-ray flux. A direct implication of this model is that the variations in power-law X-ray emission must be observed after the disc black-body variations with a time-delay scaling with the viscous travel-time between the radii where the fluctuations originate and the power-law emitting region.  So far, evidence for disc-driven variability has been seen in {\\it XMM-Newton} EPIC-pn observations of three hard state sources: GX~339-4, Cyg~X-1 and SWIFT~J1753.5--0127 \\citep{uttley2011}.  These data were obtained during relatively long-lived hard states at moderate luminosities, around 1~per~cent of the Eddington limit (with some uncertainty given the uncertain distances to GX~339-4 and SWIFT~J1753.5--0127 as well as some uncertainty on their masses).  More luminous hard states, corresponding to outburst rises or `failed' state transitions, have previously not been studied with detailed spectral-timing into the soft X-ray band, but could contain important information about changes in the disc variability as the accretion rate increases and the disc emission strengthens. In this paper, we carry out the first detailed spectral-timing study of a luminous hard state of SWIFT~J1753.5--0127, a transient black hole candidate which was discovered on 2005 May 30 \\citep{palmer2005}. Its $\\sim 3.2$ hr period derived from optical lightcurves \\citep{zurita2008,durant2009} makes it the BHXRB with the second-shortest period known to date \\citep{kuulkers2012}. Since the start of the outburst, it has undergone a transition to and from the hard-intermediate state. It has never completed the outburst cycle toward softer states nor has it gone into quiescence. The long-term {\\it Swift/BAT} lightcurve of SWIFT~J1753.5--0127 is shown in Figure~\\ref{fig:outb} and shows the epochs of the two {\\it XMM-Newton} EPIC-pn observations.  The first, obtained early in 2006, corresponds to a relatively faint hard state and shows clear evidence for disc-driven variability \\citep{wilkinson2009,uttley2011}.  The second observation was triggered by us in response to a brightening of the hard state in September 2009, and enables a comparison of a more luminous hard state with the lower-luminosity states studied to date.  In Section~\\ref{sec:datared} we describe our data reduction, and in Section~\\ref{sec:results} we carry out a detailed spectral-timing study of the 2009 data and compare our results with the 2006 observation.  In Section~\\ref{sec:discussion} we summarise our key results and interpret them in terms of the disc fluctuation model which can explain the data obtained at lower luminosities. ",
        "conclusions": "The soft X-ray coverage of the EPIC-pn camera onboard \\textit{XMM-Newton} has allowed us to carry out a comprehensive study of the spectral-timing properties of the BHXRB SWIFT~J1753.5--0127 in the luminous 2009 hard state, and to perform a comparison with the fainter 2006 hard state.  In our previous study of a sample of hard-state BHXRBs \\citep{uttley2011}, we showed that (i) soft (disc) X-ray variations lead hard (power-law) variations up to frequencies of about 1 Hz, and (ii) at the same low frequencies, the fractional variability in the disc is larger than the variability observed in the power-law component. Mass-accretion rate fluctuations that propagate through the disc would produce variability on frequencies corresponding to the viscous time-scale where they are generated, and eventually reach a centrally-concentrated source of power-law photons. This basic picture would both explain the hard-to-soft lags and give a reasonable explanation for the excess fractional variability seen in the disc as compared to the power-law.  The 2009 observation shows a breadth of previously unknown observational features, which we summarise as follows.  Firstly, the variability amplitudes in the PSD are strongly dependent on energy, with the harder band varying more than the softer bands. The explanation for this effect is seen in the fractional covariance spectrum, where the bright disc black-body component that is seen in the spectrum now appears suppressed in fractional-variability terms with respect to the power-law, at all frequencies considered. This is opposite to what was found in less-luminous hard states in \\citet{uttley2011}, where the disc was more variable than the power-law component. The higher luminosity (and thus, accretion rate) seen in this observation could imply that the disc emissivity in the observed bands is less centrally peaked than in less-luminous hard states. In this case, a fluctuation that is produced at an inner disc radius (e.g. within 1.5 $r_{\\rm in}$) will not modulate the emission from outer radii and this will effectively reduce the \\emph{observed} fractional variability of the disc as a whole. This is consistent with the shape of the energy-dependent lags.  Secondly, we have found that the low frequency ($\\lesssim 0.5$~Hz) disc variability below $< 0.6$~keV appears partly uncorrelated with the rest of the emission produced by both the disc and power-law components. An explanation for this could exist if this incoherent accretion variability is produced at larger radii than the coherent variability, affecting the softest emission that is more extended while being unable to propagate to the innermost part of the disc and the power-law emitting region due to viscous damping.  Finally, the last feature to highlight is an up-turn in the frequency-resolved time lags above $\\sim 0.5$~Hz. This could be produced if, e.g. a relatively central secondary extended corona that sandwiches the disc -- and thus reduces the observed hard-to-soft lags caused by disc propagation variations -- becomes swamped by a coherent source of variations associated with the high-frequency peak in the PSD (perhaps an even smaller unstable radius in the disc).  By identifying a number of spectral-timing signatures with the accretion disc, this work suggests that a wealth of information regarding the properties of accretion flows can be discovered by extending spectral-timing measurements to softer X-rays and using techniques that allow better spectral-resolution. By combining these powerful techniques with more detailed modelling, a comprehensive understanding of the variability may help to unravel many of the mysteries underlying accretion physics and associated phenomena such as the nature of state transitions and the production of the jets seen in the hard states of accreting black holes."
    },
    "1208/1208.0006_arXiv.txt": {
        "abstract": "Tangential discontinuities, seen as X-ray edges known as cold fronts (CFs), are ubiquitous in cool-core galaxy clusters. We analyze all 17 deprojected CF thermal profiles found in the literature, including three new CFs we tentatively identify (in clusters A2204 and 2A0335). We discover small but significant thermal pressure drops below all nonmerger CFs, and argue that they arise from strong magnetic fields below and parallel to the discontinuity, carrying $10\\%-20\\%$ of the pressure. Such magnetization can stabilize the CFs, and explain the CF---radio minihalo connection. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.2003_arXiv.txt": {
        "abstract": "The Galaxy appears to be richer in young, massive stellar clusters than previously known, due to advances in infrared surveys which have uncovered deeply embedded regions of star formation. Young, massive clusters can significantly impact the surrounding interstellar medium (ISM) and hence radio observations can also be an important tracer of their activity. Several hundred cluster candidates are now known by examining survey data. Here we report on multiwavelength observations of six of these candidates in the Galaxy. We carried out 4.9 and 8.5 GHz VLA observations of the radio emission associated with these clusters to obtain the physical characteristics of the surrounding gas, including the Lyman continuum photon flux and ionized gas mass. {\\it Spitzer} Infrared Array Camera observations were also made of these regions, and provide details on the stellar population as well as the dust continuum and polycyclic aromatic hydrocarbon emission. When compared to the known young, massive clusters in the Galaxy, the six cluster candidates have less powerful Lyman ionizing fluxes and ionize less of the H II mass in the surrounding ISM. Therefore, these cluster candidates appear to be more consistent with intermediate-mass clusters (10$^3$-10$^4$ M$_\\sun$). ",
        "introduction": "Massive stars return a significant fraction of their mass to the interstellar medium (ISM) by means of stellar winds and supernovae. Massive stars are rare, have short lifetimes, and are predominantly observed in young, massive ($\\le$ 20 Myr, $\\ge$ 10$^4$ M$_\\sun$) clusters. To date, few such clusters are known and well-studied in the Galaxy because many are located in the highly obscured Galactic disk. A complete census of young, massive stellar clusters in the Galaxy is needed to determine their total number and distribution. Estimates have shown that the Galaxy should host close to 100 of these clusters (\\citealt{H-P08}). However, only a few such clusters are well known and studied. For example, the Arches, Quintuplet, and Central clusters in the Galactic center were uncovered with infrared spectroscopic observations, and are responsible for heating and shaping much of the surrounding ISM (\\citealt{Cotera96}; \\citealt{Figer99a}b; \\citealt{Lang97}, \\citeyear{Lang01a}, \\citeyear{Lang02}). In addition, several other young, massive clusters in the Galaxy have been recently studied: RSGC 1 \\&~2 (\\citealt{Davies08}), W49a (\\citealt{H-A05,de Pree97}), NGC 3603 (\\citealt{Melena08,Borissova08,de Pree99}), and Westerlund 1 (\\citealt{Brandner08,K-D07}).  \\paragraph{}Over the last decade, surveys in the near-infrared and mid-infrared (2MASS and {\\it Spitzer}/GLIMPSE in particular) have revealed as many as 1000 candidate clusters (e.g. \\citealt{Bica03,Mercer05}). A sample of 40 possible young, massive clusters were selected from these candidates using multiwavelength data, including radio and optical observations. Among these 40 sources, six were identified as having bright associated radio emission from the NRAO VLA Sky Survey (VLA-NVSS; \\citealt{Condon98}). Here, we present a multiwavelength study of these six cluster candidates in the Galaxy. The radio observations in this study have higher sensitivity and resolution than the VLA-NVSS data, and provide greater insight into the properties of the gas and dust surrounding these clusters. The primary goal of the radio data is to determine the properties of the ISM surrounding the candidate massive clusters, and to compare them to the known young, massive clusters in the Galaxy. Estimates of the ionized mass and ionizing flux can be obtained with the radio data, from which the candidate sources can be categorized as supermassive or intermediate-mass clusters. The primary goal of the infrared data is to characterize the infrared emission surrounding the clusters as well as the infrared emission arising from cluster members. Comparison between the radio and infrared observations provides additional insight about the interstellar environment of the stellar clusters and the impact of the stars on the surrounding medium. ",
        "conclusions": "Observations have been carried out using {\\it Spitzer}/IRAC at 3.6, 4.5, 5.8 and 8.0 $\\mu$m and VLA radio observations at 8.5 and 4.9 GHz to determine the properties of six candidate stellar clusters and their surrounding H II regions. The following conclusions have been drawn from this study:  \\begin{enumerate} \\item Image overlays of all the sources were created to compare the distribution of infrared and radio emission. The 3.6 and 8 $\\mu$m emission was found to be more diffuse than the radio emission in the majority of the cluster sources. In addition, the peak intensity in the radio emission is often offset from the peak intensity in the infrared emission. \\item Radio continuum properties of the H II regions led to an estimate of the sources' ionized mass and Lyman continuum photon flux. BD65 is the most massive of the candidate clusters having an ionized mass of $\\sim$2800 M$_\\sun$ and N$_{Lyc}$ $\\sim$10$^{49}$. DBCL23 was also found to have a fairly large ionizing flux of $\\sim$10$^{48}$. The physical parameters of the sources from the radio continuum were compared to H II parameters of other well-studied young, massive Galactic clusters. \\item A more complete multiwavelength analysis of the morphology in BD52, DBCL23, and BD84 was carried out. The emission mechanisms in the clusters and surrounding ISM could be explored by investigating the intensity of emission at different wavelengths across a slice through the sources. A general picture emerged from this analysis: ionized gas surrounded by bright warm dust and PAH emission. This is consistent with other studies of infrared dust bubbles (\\citealt{Watson08}). \\item Infrared photometry of stellar members of the candidate clusters was used to search for candidate YSOs. The existence of YSOs provides evidence of ongoing star formation. Candidate YSOs were found in all the cluster candidates using color-color plots. A lack of detection of MYSOs is consistent with the general picture for all six clusters: young, partially embedded, intermediate mass stellar clusters. \\item Comparison with other young, massive clusters (e.g. Arches, Quintuplet, Westerlund 1) leads to the conclusion that the stellar cluster candidates presented in this study are not massive ( $>$ 10$^4$ M$_\\sun$) clusters. More likely, they represent examples of Galactic intermediate mass ( $>$ 10$^3$ M$_\\sun$) stellar clusters. \\end{enumerate}"
    },
    "1208/1208.0230_arXiv.txt": {
        "abstract": "We report on sensitive dual-frequency (1.7 and 5 GHz) European VLBI Network observations of the central region of nine Seyfert galaxies. These sources are among the faintest and least luminous members of a complete sample of nearby ($d<22$ Mpc) low luminosity AGNs. We detect radio emission on milliarcsecond scale in the nuclei of 4 galaxies, while for the other five sources we set an upper limit of $<\\sim100\\ \\mu$Jy. In three sources, namely NGC\\,3227, NGC\\,3982, and NGC\\,4138, radio emission is detected at both 1.7 and 5 GHz and it is resolved in two or more components. We describe the structural and spectral properties of these features; we find that in each of these three nuclei there is one component with high brightness temperature (typically $T_B >10^{7.5}$~K) and flat/intermediate spectral index ($0.3\\le \\alpha \\le 0.6$, $S(\\nu)\\sim\\nu^{-\\alpha}$), accompanied by secondary steep spectrum extended components. In these cases, non-thermal emission from jets or outflows is thus the most natural explanation. A faint feature is detected in NGC\\,4477 at 5 GHz; keeping in mind the modest significance of this detection ($\\sim 5\\sigma$), we propose the hot corona as the origin of non-thermal emission, on the basis of the unrealistic magnetic field values required by synchrotron self-absorption. Finally, the five non-detected nuclei remain elusive and further observations on intermediate scales will be necessary to investigate their nature. ",
        "introduction": "Active galactic nuclei are traditionally divided into radio quiet (RQ) and radio loud (RL) depending on the ratio between optical and radio flux density \\citep{Kellerman1989}. The nuclear regions of RL AGNs are characterized by high brightness temperature cores and jets with compact knots, in some cases moving with superluminal velocity; on large scale, the radio emission of RL AGNs can reach out to several hundred kiloparsecs, well beyond the size of the host galaxy. RQ AGNs like Seyfert galaxies are much fainter in the radio band and their emission is confined in the sub kpc scale. However, it has become clear in the last years, through VLA surveys, that RQ AGNs are not completely silent in the radio band \\citep[e.g.][]{Nagar2002,Ho2001}. While the origin of the radio emission is clear in RL AGNs as synchrotron radiation from energetic particles in jets and lobes, the case of RQ AGNs is not well established. Since the radio structure in the nuclear region is complex, it is of fundamental importance to resolve them with the highest spatial resolution achievable and to obtain spectral information through multi-wavelength radio observations. In some bright targets, VLBI studies have successfully shed light on the properties of these regions, e.g.\\ revealing thermal free-free emission from an X-ray heated corona in NGC\\,1068 \\citep{Gallimore2004}, or two-sided jet-like structures with low speeds, indicating non-relativistic jet motion, possibly due to thermal plasma like in NGC\\,4151 \\citep{Ulvestad2005} or at most mildly relativistic motion of non-thermal plasma as in NGC\\,4278 \\citep{Giroletti2005}.  However, the bulk of the low luminosity radio quiet AGN population is characterized by low radio flux densities which require high sensitivity for a proper study. We thus considered the complete and distance limited sample of 28 local ($d\\le22$ Mpc) Seyfert galaxies selected by \\citet{Cappi2006}\\footnote{The original paper only counted 27 sources, since NGC\\,3982 lacked {\\it XMM-Newton} data at that time}. In \\citet{Giroletti2009}, we observed the cores of five weak ($\\sim$ mJy level) targets with the European VLBI Network (EVN) to complement observations of the brighter members of the sample available in the literature \\citep[e.g.][]{Trotter1998,Gallimore2004,Ulvestad2005} . Of these, 4 sources were detected, revealing a complex scenario, where diverse underlying physical mechanisms can be responsible for the nuclear radio emission in the four detected targets. The radio spectral indices $\\alpha$ (defined such that $S(\\nu)\\sim\\nu^{-\\alpha}$) range from steep ($\\alpha>0.7$) to slightly inverted ($\\alpha=-0.1$), brightness  temperature vary from $T_B=10^5$ K to larger than $10^7$ K and cores are either resolved or unresolved, in one case (NGC\\,4051) accompanied by a lobe-like structure  \\citep{Giroletti2009}.  In this paper, we continue the study of the sources in the above mentioned sample discussing dual-frequency EVN observations for 9 more sources never observed on milliarcsecond scale. With this work, we complete the VLBI observations for all the sources in the sample that have a detection on VLA scales in at least one radio frequency. In a companion paper (Panessa et al.\\ in prep.), we attempt a statistical study of the multi-wavelength properties of the whole sample.  This paper is structured as follows:  in Sect.~\\ref{s.observations} we describe our observations and the data reduction procedures;  the results are presented in Sect.~\\ref{s.results} and their discussion is given in Sect.~\\ref{s.discussion}; finally, we summarize the main conclusions in Sect.~\\ref{s.conclusions}. ",
        "conclusions": "The aim of this work was to observe the nuclear region of 9 local Seyfert galaxies in the radio band on the parsec scale. These galaxies have been extracted from the complete sample of 28 Seyfert studied in \\citet{Cappi2006} with a VLA detection but still lacking  high resolution (VLBI) observations. At 1.7 GHz we revealed 3 sources (NGC\\,3227, NGC\\,3982, and NGC\\,4138), while at 5 GHz we were able to reveal radio emission also in the core of NGC\\,4477, raising the detection rate to $\\sim 44\\%$.  In the sources detected at both frequencies, i.e.\\ NGC\\,3227, NGC\\,3982, and NGC\\,4138, we detected one component with high brightness temperature ($\\log T_B>7.5$) and flat/intermediate spectral index ($0.3\\le \\alpha \\le 0.6$), which we ascribe to non-thermal emission from the immediate vicinity of the central black hole; moreover, steep spectrum extended components are also detected within some tens to hundred milliarcseconds from the core, suggesting the presence of jets or outflows on parsec scales. Indeed, the VLBI structure in NGC\\,3227 connects nicely to the larger scale emission observed in literature MERLIN images. The physical parameters estimated under the assumption of minimum energy are quite reasonable; e.g., the equipartition magnetic field are of a few mG, and slightly higher in the most compact features.  In NGC\\,4477, which is detected only at 5 GHz, the brightness temperature is somewhat lower, and the physical parameters seem at odd with a synchrotron self-absorption scenario, mainly because of the too high magnetic field required. On the other hand, a thermal free-free origin for its radio emission seems more viable, similar to what found in NGC\\,1068 by \\citet{Gallimore2004}. However, we remind that the significance of the detection is limited ($\\sim5\\sigma$) and future observations will be needed to confirm this speculation.  Finally, NGC\\,3185, NGC\\,3941, NGC\\,4639, NGC\\,4698, and NGC\\,5194 remain undected down to very low brightness levels ($<\\sim100\\ \\mu$Jy beam$^{-1}$). Since all these sources have weak but compact cores in VLA images, observations at intermediate resolution will be necessary to reveal the structure and characterize the physical condition of their nuclear regions."
    },
    "1208/1208.2690_arXiv.txt": {
        "abstract": "\\input{abstract} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.3375_arXiv.txt": {
        "abstract": "{The Sloan Digital Sky Survey (SDSS) allows us to classify galaxies using optical low-ionization emission-line diagnostic diagrams. A cross-correlation of the SDSS data release 7 (DR7), containing spectroscopic data, with the Very Large Array (VLA) survey Faint Images of the Radio Sky at Twenty-centimeters (FIRST), makes it possible to conduct a joined multiwavelength statistical study of radio-optical galaxy properties on a very large number of sources.}{Our goal is to improve the study of the combined radio-optical data by investigating whether there is a correlation between the radio luminosity at $20$ cm over the luminosity of the optical H$\\alpha$ line ({$L_{20~cm}/L_{H\\alpha}$) and line excitation ratios, where the latter provide the spectroscopic classification in Seyferts, low-ionization nuclear emission-line regions (LINERs), and star-forming galaxies. We search for a trend with $z$ in the classification provided by classical and more recent optical emission-line diagnostic diagrams}.}{We cross-matched the optical sources with the FIRST radio survey in order to obtain spectroscopic information of a selected sample of radio emitters with optical observed counterpart. We searched for an $L_{20~cm}/L_{H\\alpha}$ threshold value above which the radio emitters start being classified as active galactic nuclei (AGNs) rather than star-forming galaxies (SFGs). We investigated the origin of emission-lines by using both photoionization and shock models.}{The percentage of detected AGNs (Seyferts and LINERs) or composites is much higher in the optical-radio sample than in the optical sample alone. We find a progressive shift in the sources towards the AGN region of the diagram with increasing $L_{20~cm}/L_{H\\alpha}$, with an indication of different behavior for LINERs and Seyferts. The classification appears to slightly depend on the redshift. The diagnostic diagrams display a density peak in the star-forming or composite region for log$(L_{20~cm}/L_{H\\alpha})<0.716$, while the distribution in the LINER region peaks above this threshold. A comparison with photoionization and shock models shows that the large fraction of LINERs identified in our study have emission lines that may be explained by shocks.}{Our results indicate that it is worthwhile to further explore the radio domain, probing the physical nature of LINERs, thanks to a combination of optical and radio information. The {[N\\sc ii]}/H$\\alpha$ vs. equivalent width of the H$\\alpha$ line (WHAN) diagram confirms the LINER classification for most of those that have been identified with the traditional diagnostic diagrams. The correlation between $L_{20~cm}/L_{H\\alpha}$ and optical emission line ratios suggests the nuclear origin of the emission from the most powerful radio galaxies.} ",
        "introduction": "Our knowledge of global galaxy properties has been recently improved thanks to large-area surveys, such as the Sloan Digital Sky Survey \\citep[SDSS;][]{York2000}, which have allowed for statistical studies on large samples of galaxies. SDSS data can be used as a starting point to conduct multiwavelength studies on galaxy properties outside the optical regime, such as the radio domain, where active galactic nuclei (AGNs) can be detected.\\newline Radio galaxies are active galaxies that are very luminous at radio wavelengths. They are considered as radio-loud if they have ratios of radio (at $5$ GHz) to optical (B-band) flux greater than ten \\citep{Kellermann1989}. According to \\citet{Fanaroff1974}, radio-loud galaxies can be divided into two major classes. In Fanaroff-Riley Class I (FR I) sources, the radio emission peaks near the galaxy nucleus and the emission from the jets fades with distance from the center, while class II (FR II) sources present bright radio lobes. Class I sources dominate the population of radio emitters at low radio power and low redshifts, while more powerful radio galaxies % are almost exclusively FR II systems that can be detected at higher redshift.\\newline Radio galaxies can also be classified according to whether they have strong high-excitation narrow-line emission. The majority of FR I radio galaxies show very weak or completely absent emission lines. Those are referred to as low-excitation systems \\citep{Hardcastle2006}, and they are mostly found in elliptical galaxies with little ongoing star formation \\citep{Ledlow1995,Govoni2000}. If optical spectroscopic information is available, these galaxies are generally classified as low-ionization nuclear emission-line regions \\citep[LINERs,][]{Heckman1980}. Conversely, the most powerful, high-redshift FR II radio galaxies have, in most cases, strong emission lines and are classified as AGNs \\citep{Wierzbowska2011} with peculiar optical morphologies, e.g. tails, bridges, and shells \\citep{Smith1989} and bluer colors with respect to giant ellipticals. It has also been established that powerful FR IIs show a strong correlation between their radio luminosity and their optical emission-line luminosity \\citep{Baum1995}, suggesting that both the optical and the radio emission originate in the same physical process.\\newline The possibility of a more general correlation between the emission lines and radio luminosities of radio AGN has been already explored in the past \\citep{McCarthy1993,Zirbel1995}. Radio-loud AGNs are most likely to display emission lines in galaxies with low velocity dispersions and at radio luminosities greater than $10^{25}~WHz^{-1}$ \\citep{Kauffmann2008}. A similar correlation has been observed between the emission-line luminosity and the ionization state of the gas for a sample of low-$z$ radio galaxies \\citep{Saunders1989}, where the higher values of the emission-line luminosities have been measured for the more powerful radio sources, as an indication of the presence of a strong ionizing AGN-like field. However, a problem arises since the most powerful radio galaxies are generally detected at higher redshifts than the less powerful radio galaxies, due to selection effects. This makes it difficult to establish whether the correlation is between the emission-line luminosity and radio luminosity, rather than between emission-line luminosity and redshift \\citep{McCarthy1993}. The well-known effect of having an increasing number of detected powerful radio AGNs with increasing redshift is not only due to a selection bias, but is is also supported by the downsizing scenario of galaxy evolution \\citep[e.g.][]{Thomas2005,Cimatti2006}. Galaxies placed at higher redshift are more massive and host massive black holes that accrete producing powerful jets. These are easily detected in the radio domain, while low-redshift radio galaxies host low-mass black holes, in which only weak jets originate. The quasars detected in the radio but not in the optical regime might be heavily obscured objects. They indeed experience dust obscuration when the line of sight passes through the optically thick torus, which surrounds the black hole (unified model, Antonucci 1993).\\newline One of the current scenarios of galaxy evolution includes the possibility of a smooth transition from late-type galaxies (LTGs), which show blue colors and ongoing star formation, to red and passively evolving early-type galaxies (ETGs). This transition could be driven by the so-called AGN-feedback \\citep{Cattaneo2007,Cattaneo2009}, which is thought to be responsible for star formation quenching. After the quenching, a phase of passive evolution starts, and the galaxies color turns into red. Some studies show that there is indeed a correlation between the age of galaxy stellar populations and the AGN activity \\citep[][Vitale et al. 2012, in prep.]{Tortora2009}, % with the oldest stars inhabiting the AGN-like galaxies. Moreover, AGNs are found to reside almost exclusively in massive galaxies with structural properties similar to normal early-type systems \\citep{Heckman2004}, which are dominated by old stellar populations.\\newline The question of whether the AGN-feedback is also capable of triggering star formation in the vicinity of the black hole still has to be answered. The standard accretion mode of AGNs, which is associated with quasar activity \\citep{Shakura1973}, is related to star formation in the host galaxies \\citep{Kauffmann2003a}. This scenario seems to contrast with the quenching of the starburst activity due to AGN-feedback. Star formation and accretion around the black hole, both fueled by the same material, might occur with delays with respect to each other \\citep{Wild2010,Tadhunter2011} and come to an end once the gas is exhausted. A fraction of those AGNs that accrete in `quasar\u00b4 mode show powerful radio jets and are therefore radio-loud.\\newline As an observational proof of this scenario, \\citet{Ivezic2002} have shown that a sample of optically unresolved radio sources from the FIRST survey has bluer colors than do other SDSS objects, and Richards et al. (2002) find that the SDSS quasar candidates - which are likely to have a radio counterpart - display blue colors, especially at $z>1$, indicative of ongoing star formation. To understand the importance of the AGN-feedback in the current scenarios of galaxy evolution, a study of the AGN populations, the properties of the host galaxies, and the physical mechanisms that trigger the production of the emission lines or the radio-activity, is necessary.\\newline AGNs can be selected from a spectroscopic survey using some optical emission line ratios. Emission-line diagnostic diagrams have been extensively used during past decades to point out the connection between the galaxy nuclear activity, its morphological type \\citep{Ho1997}, and its evolutionary stage \\citep{Hopkins2006}. The Baldwin-Phillips-Terlevich (BPT) diagnostic diagram \\citep{Baldwin1981} and its subsequent versions \\citep{Veilleux1987,Kewley2003,Lamareille2004,Dopita2002,Groves2004,Groves2004b} make use of emission line ratios whose strength is a function of the hardness of the ionizing field of the galaxy, the ionization parameter $U$, and the metallicity (see Sect. 3.4). Higher ratios are thought to mainly be the product of the ionization that arises due to accretion around the black hole (which implies the AGN at the center of the galaxy is in its active phase) rather than photoionization by hot massive OB stars. This diagnostic technique, largely used in the optical wavelength regime, allows differentiation of galaxies that show activity in their nuclei and starbursts.\\newline Some attempts to link optical and radio properties of a large sample of galaxies by using combined spectroscopic and photometric information have been already made in the past. Obric et al. (2006) find strong correlations between the fraction of detected AGNs at other wavelengths and optical properties such as flux, colors, and emission-line strengths. \\citet{Ivezic2002} discuss the optical and radio properties of $\\sim30\\ 000$ FIRST \\citep{Becker1995} sources positionally associated with a SDSS source by analyzing their colors, while \\citet{Best2005} compare the optical survey to both FIRST and NVSS radio surveys in order to derive the local radio luminosity functions of radio-loud AGNs and star-forming galaxies. In the SDSS Early Data Release \\citep{Ivezic2002}, $\\sim70$\\% of FIRST sources do not have an optical counterpart within $3\\arcsec$. This is probably because the majority of unmatched FIRST sources, detected down to $1$ mJy sensitivity, are too optically faint to be detected in the SDSS images. Moreover, the fraction of quasars in the FIRST catalog seem to be a strong function of the radio flux, monotonically decreasing from bright radio sources towards the FIRST radio sensitivity limit.\\newline SDSS spectra have been used to compute the line strength of several emission lines to classify galaxies into starbursts and AGNs by using diagnostic diagrams. In \\citet{Ivezic2002}, the number of radio galaxies classified as AGNs rather than starbursts ($\\sim30\\%$) is six times more than the corresponding number for all the SDSS galaxies. Furthermore, the radio emission from AGNs turns out to be more concentrated than the radio emission from starburst galaxies, as we expect from the nuclear origin of AGN emission. Radio emission is point-like in compact quasars detected at high redshift, while local galaxies tend to have larger radio sizes. This suggests that a significant amount of the radio emission either originates outside the nuclear region or that the radio lobes are resolved. Depending on the beam size, it can happen that not all the light from the galaxy is taken into account. The relative number of AGNs decreases with the radio flux, and this is consistent with the differences in the radio luminosity functions of starburst and AGNs \\citep{Machalski2000,Sadler2002}.\\newline In a multiwavelength approach, we made use of both optical - for the part concerning emission lines and diagnostic diagrams - and radio data to conduct a statistical study of the prospects of identifying radio galaxies in some well-defined regions of the low-ionization emission lines diagnostic diagrams by using a combination of radio and optical properties. A comparison of the spectroscopic measurements with some photoionization and shock models is then presented to shed light on the origin of the emission lines in AGNs and star-forming galaxies.\\newline The paper is organized as follows: In Sect. 2 we present the optical and radio samples and their cross-matching. We describe the way we obtain the final sample and its completeness. In Sect. 3 we present our results. We exploited the diagnostic diagrams for the cross-matched sample, looking for a trend with $L_{20~cm}/L_{H\\alpha}$ and redshift. We compare the spectroscopic measurements with photoionization and shock models. Discussion on the results is provided in Sect. 4. Main findings and conclusions are in Sect. 5. ",
        "conclusions": "Our main findings and conclusions follow: \\begin{itemize} \\item The requirement of detected radio emission predominantly selects active galaxies, a considerable number of AGNs and metal-rich starburst galaxies where the radio-emission mainly has a stellar origin. Radio emission correlates with optical emission-line ratios, so a cross-match allows us to classify and study radio emitters according to their optical spectral identification. \\item Emission-line diagnostic diagrams show that many radio emitters are classified as AGNs with increasing $L_{20~cm}/L_{H\\alpha}$. In particular, the increase in $L_{20~cm}$ selects powerful radio emitters such as Seyferts, while the decrease in $L_{H\\alpha}$ strongly contributes to the selection of a considerable fraction of galaxies with LINER-like emission. \\item The higher number of identified AGNs with increasing values of $z$ can indicate a true correlation between these quantities, as well as a bias, related to the selection of the most powerful radio emitters. \\item Emission-line diagnostic diagrams are useful tool for classifying sources according to the main ionizing mechanism producing emission lines, though the classification depends on the choice of the chemical species that is used in the diagram. In general, we found good agreement between the classifications given by the classical diagnostic diagrams and the {[N\\sc ii]}/H$\\alpha$ vs. equivalent width of the H$\\alpha$ line (WHAN) diagram. Diagnostic diagrams can only give hints on the nature of the observed emission-lines, making a complementary comparison between data and models necessary. \\item While the star-forming sequence in the diagnostic diagrams can be successfully fit by photoionization models, the AGN region seems to collect objects whose observed emission lines are due to different processes. Most of the radio emitters of our sample, which are mainly classified as LINERs at high  $L_{20~cm}/L_{H\\alpha}$ values, have emission lines whose ratio can be explained by fast-shock and dusty-AGN models. Shocks, closely linked to the presence of radio emission, can be produced both by stars and AGNs, so that unveiling the nature of LINERs requires a more detailed study. \\item By using diagnostic diagrams, it is possible to select populations of LINER-like objects and further distinguish between the `true-AGNs\u00b4 and the galaxies whose emission is produced by old stars. For our radio-selected sample, LINERs are more like true AGNs than retired galaxies. Resolving the spatial radio structure of these groups of LINER-like objects, as well as Seyferts that are particularly bright in the radio, would help us shed light on the physical mechanisms that are responsible for the observed radio luminosity and optical emission-line ratios. \\end{itemize}"
    },
    "1208/1208.1370_arXiv.txt": {
        "abstract": "{In some galaxy clusters powerful AGN have blown bubbles with cluster scale extent into the ambient medium. The main pressure support of these bubbles is not known to date, but cosmic rays are a viable possibility. For such a scenario copious gamma-ray emission is expected as a tracer of cosmic rays from these systems.} {Hydra~A, the closest galaxy cluster hosting a cluster scale AGN outburst, located at a redshift of 0.0538, is investigated for being a gamma-ray emitter with the High Energy Stereoscopic System (H.E.S.S.) array and the \\emph{Fermi} Large Area Telescope (\\emph{Fermi}-LAT).} {Data obtained in 20.2 hours of dedicated H.E.S.S. observations and 38 months of \\emph{Fermi}-LAT data, gathered by its usual all-sky scanning mode, have been analyzed to search for a gamma-ray signal.} {No signal has been found in either data set. Upper limits on the gamma-ray flux are derived and are compared to models. These are the first limits on gamma-ray emission ever presented for galaxy clusters hosting cluster scale AGN outbursts.} {The non-detection of Hydra~A in gamma-rays has important implications on the particle populations and physical conditions inside the bubbles in this system. For the case of bubbles mainly supported by hadronic cosmic rays, the most favorable scenario, that involves full mixing between cosmic rays and embedding medium, can be excluded. However, hadronic cosmic rays still remain a viable pressure support agent to sustain the bubbles against the thermal pressure of the ambient medium. The largest population of highly-energetic electrons which are relevant for inverse-Compton gamma-ray production is found in the youngest inner lobes of Hydra~A. The limit on the inverse-Compton gamma-ray flux excludes a magnetic field below half of the equipartition value of 16~$\\mu$G in the inner lobes.} ",
        "introduction": "At the center of some galaxy clusters powerful Active Galactic Nuclei (AGN) reside and the feedback of outbursts generated by these AGN on the embedding Intra-Cluster Medium (ICM) can be seen in several systems \\cite[for a review, see][]{mcnamara2007}. Typical signatures for an AGN -- ICM interaction are surface brightness depressions in the diffuse thermal X-ray emission of the cluster which are caused by cavities in the ICM. These cavities appear to be filled with non-thermal electrons which radiate in the radio band due to synchrotron emission \\citep[e.g.][]{birzan2004,dunn2006}. AGN-blown bubbles surrounded by thermal plasma offer the exciting possibility to constrain the energetics of these outbursts. This can be done by estimating the work that is necessary to expand the bubbles against the thermal pressure of the embedding ICM (pV work in the following). The energetics involved in this AGN activity can be enormous, in some cases even exceeding 10$^{61}$~erg \\citep{mcnamara2007}. The most powerful AGN outbursts known to date are found in MS~0735+7421 \\citep{mcnamara2005}, Hercules~A \\citep{nulsen2005a} and Hydra~A \\citep{nulsen2005b}. The AGN created bubbles in these systems have ages of about 10$^8$ years and exhibit sizes on the scale of the galaxy cluster itself.  The nature of the main pressure support agent which fills the bubbles in the ICM is not known to date. Viable possibilities for the pressure support in such systems would be relativistic particles such as hadronic cosmic rays or electrons \\citep[e.g.][]{dunn2004,ostrowski2001,hinton2007}, magnetic fields \\citep[e.g.][]{dunn2004} or hot plasma \\citep[e.g.][]{gitti2007}. The energy required to expand bubbles with volume $V$ into a surrounding ICM with pressure $p$ ranges from $2pV$ for magnetic fields to $4pV$ for relativistic fluids such as cosmic rays \\citep[e.g.][]{wise2007}.  One inevitable consequence of bubbles filled with non-thermal particles would be the production of gamma-ray emission. For the case of hadronic cosmic rays, gamma rays are produced by inelastic collisions between the high energy particles and the thermal surrounding medium \\citep[e.g.][]{hinton2007}. In case of electrons, gamma-rays are produced by up-scattering of cosmic microwave background (CMB) and infrared extragalactic background light (EBL) photons by these electrons \\citep[e.g.][]{abdo2010a}. Large-scale leptonic gamma-ray emission connected to AGN lobes has indeed been discovered with the \\emph{Fermi} satellite from the radio galaxy Centaurus~A \\citep{abdo2010a} and potentially from NGC~6251 \\citep{takeuchi2012}. Both galaxies are not hosted by a cluster. To date, no galaxy cluster has been firmly detected in gamma rays \\citep{perkins2006,aharonian2009a,aharonian2009b,aleksic2010,ackermann2010}. However, the detection of an extended  gamma-ray signal resulting from annihilation emission from supersymmetric dark matter has been claimed for the Virgo, Fornax and Coma cluster \\citep{han2012}. Recently, NGC~1275 the central radio galaxy of the Perseus cluster has been detected in VHE gamma rays \\citep{aleksic2012a}, and with this deep exposure stringent upper limits on the emission of the Perseus cluster itself have been obtained \\citep{aleksic2012b}. Galaxy clusters hosting cluster-scale AGN outbursts appear to be promising targets for gamma-ray observations according to the extraordinary energetics inferred from the AGN -- ICM interaction seen in these systems. To date, no gamma-ray observations on galaxy clusters that host cluster-scale AGN outbursts have been presented.   The Hydra~A system (Abell~0780) at RA(J2000)~=~9$^\\mathrm{h}$18$^\\mathrm{m}$05.7$^\\mathrm{s}$ and Dec(J2000)~=~-12\\degr05\\arcmin44\\arcsec at a redshift of 0.0538 is the closest known galaxy cluster which hosts a cluster-scale AGN outburst \\citep{nulsen2005b}. It features several cavities with a total expansion work $pV$ of $4 \\times 10^{60}$~erg done on the ICM. Thus the total energy required, depending on the equation of state of the main pressure agent, is $(0.8 - 1.6) \\times 10^{61}$ erg which were deposited in the last few 10$^8$ years in the surroundings \\citep{wise2007}. Hydra~A also features low-frequency radio lobes extending to almost 4\\arcmin\\ from the cluster center \\citep{lane2004}. Shocks in the ICM surround these radio lobes \\citep{nulsen2005b} with energetics of $9 \\times 10^{60}$ erg, comparable to the expansion work done in the cavities against the thermal plasma. The central AGN outburst has also driven substantial gas dredge-up in the Hydra~A system \\citep{gitti2012}.  \\emph{Chandra} has furthermore revealed an extensive cavity system consisting of three generations of cavities with decreasing ages, which points towards a complex activity history of the system \\citep{wise2007}. Most relevant for gamma-ray production are the giant outer lobes that dominate the energetics in the Hydra~A system and the inner lobes that are expected to contain the youngest population of particles. Both possibilities will be further discussed in Sec.~\\ref{sec:discussion}.  Due to its proximity and energetics, Hydra~A is expected to feature the highest gamma-ray flux of all galaxy clusters harboring cluster-scale AGN outbursts. For the case of hadron-dominated bubbles it was inferred that the flux might be close to the detection limit of the current generation of gamma-ray instruments \\citep{hinton2007}.  In this paper, upper limits on the gamma-ray emission from the Hydra~A system are reported. Limits obtained by the High Energy Stereoscopic System (H.E.S.S.) and the \\emph{Fermi} Large Area Telescope (\\emph{Fermi}-LAT) are presented in Sec.~\\ref{sec:hess} and Sec.~\\ref{sec:fermi}, respectively. These limits are used to obtain constraints on the energy of hadronic cosmic rays (Sec.~\\ref{sec:hadronic}) and electrons (Sec.~\\ref{sec:leptonic} and \\ref{sec:sync-ic}) which may populate the AGN outburst region in this galaxy cluster.   Throughout this paper a $\\Lambda$CDM cosmology with $H_0$~=~70~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\Lambda}$ = 0.7 and  $\\Omega_{M}$ = 0.3 is assumed, corresponding to a luminosity distance of $d_\\mathrm{L}$ = 240 Mpc, an angular diameter distance of 216 Mpc and a linear scale of 1.05 kpc per arcsecond \\citep{wise2007}. ",
        "conclusions": "In this paper the nearby galaxy cluster Hydra~A that hosts a cluster-scale AGN outburst is investigated for being a gamma-ray emitter. Galaxy clusters hosting a cluster-scale AGN outburst are potentially detectable gamma-ray sources due to the enormous energetics inferred from the observed AGN - ICM interactions. However, only upper limits could be obtained from 20.2 hours of H.E.S.S. and 38 month of \\emph{Fermi}-LAT observations. The non-detection of Hydra~A in gamma-rays has important implications on the particle populations and physical conditions inside the bubbles in this system. These upper limits constrain the total energy contained in relativistic particles such as hadronic cosmic rays and electrons, which can be compared to the energy which is necessary to prevent the observed cavities in the ICM from collapsing.  Constraints on the particle population in the Hydra~A galaxy cluster can also be compared to the limits on such a non-thermal component inferred for the Perseus cluster \\citep{aleksic2012b}. The Perseus cluster is an interesting candidate for this comparison since it also shows signatures of AGN -- ICM interactions in form of radio lobes and cavities in the ICM \\citep[e.g.][]{fabian2011}. For the Perseus cluster \\citet{aleksic2012b} constrained the average fraction of energy in hadronic cosmic rays to thermal energy $E_\\mathrm{CR}/E_\\mathrm{th}$ to $\\lesssim 1-2$\\% depending on the exact assumptions. For Hydra~A the hadronic cosmic ray content in the central 200~kpc can be limited to about $5 \\times 10^{60}$\\,erg, assuming complete mixing between cosmic rays and ICM (see Fig. \\ref{figure:sed}). When adopting a total gas mass of $5 \\times 10^{12}$\\,M$_\\odot$ and a temperature of 4 keV for the central 200\\,kpc of Hydra~A \\citep{david2001} then $E_\\mathrm{CR}/E_\\mathrm{th} \\lesssim 13$\\% is found. This is significantly less constraining than for the case of the Perseus cluster. Hydra~A, however, is the prime candidate to explore the particle content of giant AGN-blown lobes in galaxy clusters. This follows from the fact that the AGN outburst in the Perseus cluster is about an order of magnitude less energetic \\citep[$pV \\approx 3 \\times 10^{59}$\\,erg,][]{fabian2011} than the AGN feedback in Hydra~A. The smaller energetics is not readily compensated by the shorter distance to the Perseus cluster ($d_\\mathrm{L}$ = 75~Mpc).  For Hydra~A for the case of bubbles mainly supported by hadronic cosmic rays these upper limits can exclude the most favorable model, that requires full mixing between relativistic particles and embedding thermal medium. It is found that \\emph{Fermi}-LAT can constrain the degree of mixing to 50\\% and H.E.S.S. can limit the degree of mixing to less than 70\\%. However, hadronic cosmic rays still remain a viable pressure support for the bubbles.  In contrast to hadrons, electrons cool quite fast above GeV energies in the environment of the Hydra~A system. Consequently, a passively evolving population of electrons in the oldest outer lobes cannot be detected with the presented observations. However, for the youngest, inner radio lobes, the limit on the IC flux seems to exclude a magnetic field below about 8 \\,$\\mu$G, that is half of the equipartition value. For the large outer lobes, a population of electrons rejuvenated by in situ particle acceleration comparable to the one detected in the Centaurus~A system can be excluded as the main pressure support of the bubbles. Upper limits in the VHE gamma-ray range are not constraining for leptonic scenarios with respect to limits obtained in the GeV range.  The main feedback agent which drives the evolution of the cavities in the ICM in the Hydra~A galaxy cluster still remains unidentified. The upcoming Cherenkov Telescope Array \\citep[CTA,][]{actis2011} with its increased sensitivity will be crucial to test especially the presence of hadronic cosmic rays in the Hydra~A galaxy cluster."
    },
    "1208/1208.1693_arXiv.txt": {
        "abstract": "Tentative line emission at 111 and 129 GeV from 16 unassociated {\\it Fermi}-LAT point sources has been reported recently by \\citet{meng}. Together with similar features seen by {\\it Fermi} in a region near the Galactic Centre, the evidence has been interpreted as the spectral signature of dark matter annihilation or internal bremsstrahlung. Through a combination of supervised machine-learning algorithms and archival multiwavelength observations we find that 14 out of the 16 unassociated sources showing the line emission in the Su \\& Finkbeiner sample are most likely active galactic nuclei (AGN). Based on this new evidence, one must widen the range of possible solutions for the 100--140 GeV excess to include a very distinct astrophysical explanation. While we cannot rule out a dark matter origin for the line emission in the Galactic Centre, we posit that if the detection in the Su \\& Finkbeiner sample is indeed real it might be related to accretion, bubble, or jet activity in nearby ($z < 0.2$) AGN. Alternatively, given the right conditions, the similarity could be due to a chance occurrence caused by extragalactic background light (EBL) absorption. Or else one must concede that the features are an artefact of instrumental or calibration issues. ",
        "introduction": "Frantic activity has ensued over the past few months following the report of an excess of {\\it Fermi} gamma-ray events clustered around 100 and 140 GeV in a region near the Galactic Centre \\citep{weniger, tempel,su1}, as well as in galaxy clusters \\citep{hektor1}. Dark matter annihilation and internal bremsstrahlung have rapidly emerged as possible explanations \\citep{bring,weniger}. Alternative interpretations have also been advanced \\citep{profumo,boyarsky, aharonian}. More recently things have heated up even further with a tantalising claim of similar line emission at 111 and 129 GeV in 16 unassociated sources detected in the Second {\\it Fermi}-LAT catalogue (2FGL). The detection could provide independent support for a dark matter origin for the line emission seen near the Galactic Centre region \\citep{meng}. Certainly, such coincidence might not only help us unlock the mysteries of dark matter, but it would also prove the existence of dark matter subhaloes \\citep{klypin,moore,springel}.  In the absence of obvious flaws in the analysis, the collected evidence has risen as a sort of dark knight -- albeit an indirect one -- that might finally grant us non-gravitational access to dark matter. Intrigued by this possibility, we explore the nature of the 16 {\\it Fermi} unassociated sources listed by \\citet{meng}. Based on machine-learning classification algorithms and multiwavelength examination, we show that 14 out of the 16 unassociated {\\it Fermi} sources displaying the lines are likely gamma-ray AGN. Therefore, rather than strengthening the argument, the detection of an identical signal in the Su \\& Finkbeiner sample appears to disprove a dark matter origin for the {\\it Fermi} features unless a set of very unique astrophysical conditions are met.  The paper is structured as follows. In Section \\ref{sibyl} we explain the machine-learning classifier {\\it Sibyl}. In Section \\ref{classp} we compile class prediction for the 16 unassociated {\\it Fermi} sources listed in \\citet{meng}. Section \\ref{multi} details multiwavelength searches for the potential counterparts of these 16 objects. Finally, we provide some interpretation in Section \\ref{interp}. ",
        "conclusions": "\\label{interp} We have presented class predictions of the Random Forest classifier {\\it Sibyl} for 16 unassociated {\\it Fermi} sources showing line emission at 111 and 129 GeV. We find that 14 out of 16 unassociated sources in the Su \\& Finkbeiner sample are AGN candidates with prediction accuracy rates greater than 97.1\\%. In addition, we have detected 10 X-ray and 13 radio potential counterparts distributed over the 16 unassociated \\fermi\\ 95\\% confidence error ellipses that would be consistent with the AGN predictions. We emphasise the word potential here as a more exhaustive detective work must be completed to confirm the appropriate counterpart for each unassociated source.  It was postulated that the gamma-ray lines among the unassociated were perhaps connected to dark matter subhaloes dragged into the Galactic disc \\citep{meng}. However, assuming an isotropic distribution, at least 160 {\\it Fermi AGN} are expected at $|b| \\leq 10^\\circ$. To date only about 100 are accounted for in the 2FGL \\citep{2lac}. Thus, it makes astrophysical sense that AGN are making up an important fraction of the Su \\& Finkbeiner sample even at relatively low Galactic latitudes.  In light of these results, the dark matter origin for the narrow gamma-ray features observed by {\\it Fermi} is in question. Were these dark matter subhaloes \\citep{baltz,diemand,kuhlen}, coincidence between the Galactic Centre and the Su \\& Finkbeiner sample would certainly confirm a dark matter particle origin \\citep{hooper}. However, the interpretation changes dramatically if the unassociated sources showing an identical line signature are AGN, as implied by both machine-learning classifiers and the multiwavelength arguments just presented. Dark matter could be fed into AGN jets and the Galactic Centre, but such an explanation feels contrived given the hadronic and leptonic dominance in the gamma-ray photon field \\citep{hinton}.  Instead, a distinct astrophysical mechanism unrelated to dark matter annihilation and linked to nearby AGN ($z < 0.2$ to avoid redshifted lines) such as accretion,  bubble \\citep{meng1,profumo}, or jet \\citep{su2} phenomenology would appear to be more logical. However, we note that although many {\\it Fermi} AGN display photons above $\\sim$ 10 GeV, only a handful of soft AGN ($\\Gamma > 2$) exhibit maximum photon energies greater than 100 GeV at $z > 0.5$ \\citep{2lac}. Consequently, \\citet{meng} might be detecting a fiendish cluster of events imprinted by EBL absorption in the same energy band, but completely unrelated in origin to the emission observed near the Galactic Centre region.  Oddly enough, the lines reported by \\citet{meng} appear to be only present collectively in unassociated sources and do not appear as pronounced among associated sources, including well-known gamma-ray AGN \\citep{meng}. Therefore, we must also admit the possibility that the spectral signatures detected by {\\it Fermi} originate from confounding instrumental or calibration problems \\citep{hooper,hektor2,hektor3,fink}. The {\\it Fermi} calibration team will have the final word on the matter very soon, but independent efforts must be made to scan the public {\\it Fermi} archive for gamma-ray lines among individual AGN at $z < 0.2$, as well as in diffuse emission outside the Galactic plane.  We shall hear more about this energy region by the end of the year with the recently unveiled H.E.S.S. II \\citep{becherini,berg}, and even more sensitive observations will be available later on after completion of the Cherenkov Telescope Array  \\citep{cta}. In the future, a dark knight might rise again. Until then, we eagerly await for the final chapter of this intriguing saga."
    },
    "1208/1208.3143_arXiv.txt": {
        "abstract": "{Red supergiant stars (RSGs) and yellow hypergiant stars (YHGs) are believed to be the high-mass counterparts of stars in the asymptotic giant branch (AGB) and early post-AGB phases. As such, they are scarcer and the properties and evolution of their envelopes are still poorly understood. } {We study the mass-loss in the post main-sequence evolution of massive stars, through the properties of their envelopes in the intermediate and warm gas layers. These are the regions where the acceleration of the gas takes place and the most recent mass-loss episodes can~be~seen.} {We used the HIFI instrument on-board the Herschel Space Observatory to observe sub-millimetre and far-infrared (FIR) transitions of CO, water, and their isotopologues in a sample of two RSGs (NML\\,Cyg and Betelgeuse) and two YHGs (IRC+10420 and AFGL\\,2343) stars. We present an inventory of the detected lines and analyse the information revealed by their spectral profiles. A comparison of the line intensity and shape in various transitions is used to qualitatively derive a picture of the envelope physical structure. On the basis of the results presented in an earlier study, we model the CO and \\treceCO\\ emission in IRC+10420 and compare it to a set of lines ranging from the millimetre to the FIR.} {Red supergiants have stronger high-excitation lines than the YHGs, indicating that they harbour dense and hot inner shells contributing to these transitions. Consequently, these high-$J$ lines in RSGs originate from acceleration layers that have not yet reached the circumstellar terminal velocity and have narrower profiles than their flat-topped lower-$J$ counterparts. The YHGs tend to lack this inner component, in line with the picture of detached, hollow envelopes derived from studies at longer wavelengths. NH$_3$ is only detected in two sources (NML\\,Cyg and IRC+10420), which are also observed to be the strongest water-line emitters of the studied sample. In contrast, OH is detected in all sources and does not seem to correlate with the water line intensities. We show that the IRC+10420 model derived solely from millimetre low-$J$ CO transitions is capable of reproducing the high-$J$ transitions when the temperature in the inner shell is simply lowered by about 30\\%.} {} ",
        "introduction": "\\label{intro}  Red supergiant stars (RSGs) and yellow hypergiant stars (YHGs) are thought to be stages in the post main-sequence evolution of stars with initial masses between $\\sim$10 solar masses and 50 \\ms\\ (e.g.\\ de Jager 1998, Meynet \\& Maeder 2003, Levesque 2010). As such, RSGs and YHGs are the massive counterparts of AGB and (early) post-AGB stars. However, the very different properties of evolved massive stellar objects (RSGs, YHGs, Wolf-Rayet stars, luminous blue variable stars, Cepheid-like variables, supernovae etc), particularly their distribution in the H-R diagram, is hard to interpret with a simple description of the evolution in this phase.  Both RSGs and YHGs are known to show complex mass-loss phenomena, which form circumstellar envelopes (CSEs) that can be very dense. Mass-loss plays an important role in the evolution of these objects (de Jager 1998, Meynet \\& Maeder 2003): almost one half of their total initial mass can be ejected during these late phases, affecting in particular their possible later evolution into supernova events. Some objects have very high mass-loss rates, with episodic rates as high as \\mloss\\ $\\sim$ 10$^{-3}$ \\my, while the circumstellar envelopes around others are very diffuse, corresponding to rates under 10$^{-7}$ \\my\\ (e.g. Castro-Carrizo et al.\\ 2007 - CC07 thereafter, Quintana-Lacaci 2008, Mauron \\& Josselin 2011). From the theoretical point of view, very high rates are also expected, at least episodically, owing to a combination of high radiation pressure and atmospheric activity (e.g.\\ Josselin \\& Plez 2007; as happens in the AGB) or to the intrinsic instabilities characteristic of the yellow phases of the late evolution of high-mass stars (the so-called \"yellow void\"; see e.g.\\ de Jager 1998).  Molecular lines have been one of the most powerful tools for the study of the properties of CSEs around AGB and post-AGB stars. Similar data on RSGs and YHGs are rare, but in some cases molecular lines have yielded a quite comprehensive insight into their CSEs. For instance, using model fitting of mm- and submm-observations of several CO rotational lines in the RSG VY CMa, Decin et al.\\ (2006) deduced mass-loss rates as high as 3$\\times$10$^{-4}$ \\my, that varied on timescales of about 1000 yr.  VY CMa is known to harbour a wide variety of molecular species, illustrating the very complex chemistry at play in its CSE (e.g.\\ Tenenbaum et al.\\ 2010 and references therein). Other less well-studied RSGs, such as VX Sgr and NML Cyg, also show high mass-loss rates of the order of $\\sim$ 10$^{-4}$ \\my (e.g.\\ De Beck et al.\\ 2010). An example of a low-\\mloss\\ RSG is Betelgeuse ($\\alpha$\\,Ori), with a value of about 10$^{-7}$-10$^{-6}$ \\my, although this value is uncertain because its molecular emission may be weaker than in other objects owing to its molecular underabundance (see Huggins et al.\\ 1994, Mauron \\& Josselin 2011).  The YHGs IRC+10420 and AFGL\\,2343 have also been well-studied in molecular emission. CC07 (see also Quintana-Lacaci et al.\\ 2008) compared line-emission models with high-resolution maps of the CO \\juc\\ and \\jdu\\ lines, deriving large \\mloss\\ variations on timescales of about 1000 yr and \\mloss\\ maxima in excess of 10$^{-3}$ \\my. These objects also have many molecular lines and a very rich chemistry (Quintana-Lacaci et al.\\ 2007). The total masses in the molecule-rich shells are particularly high for these two YHGs, $M_{\\rm tot}$ \\gsim\\ 1 \\ms. The total masses of the molecular envelopes around RSGs are somewhat lower, between 0.1 solar masses and 1 \\ms\\ for VY CMa, VX Sgr, and NML Cyg.  Despite the observational progress that has been made to date, the available data still fail to provide information on some basic parameters. Low-$J$ CO lines are good tracers of both the circumstellar mass distribution and kinematics, but these easily excited lines cannot probe warm gas with temperatures $T_{\\rm k}$ \\gsim\\ 100 K, which are expected to be present and sometimes dominant in these shells. The temperature itself is not well-determined under these conditions, which may affect the determination of the mass-loss rate and the total mass. To properly study these warm components, it is therefore necessary to observe lines in the far-infrared (FIR), involving level excitations comparable (in temperature units) to these moderate kinetic temperatures.  This paper presents {\\it Herschel}/HIFI observations in the FIR of molecular lines from two RSGs, Betelgeuse and NML Cyg, and two YHGs, IRC+10420 and AFGL\\,2343. The observations are part of the Herschel guaranteed time key program HIFISTARS (PI V. Bujarrabal), devoted to the study of high-excitation molecular lines in (low- and high-mass) evolved stars. Companion observations of VY CMa, which are particularly rich and complex, will be discussed in another paper (Alcolea et al., in prep.). We present data of the \\jsc, \\jdn, and \\jdsq\\ lines of \\doce\\ and \\trece, which are good probes of the excitation. We also discuss our observations of other molecules, including water vapour, OH, and NH$_3$, which are particularly useful to understanding more clearly the chemistry in these unusual objects. ",
        "conclusions": "\\label{conclusion}  We have reported {\\it Herschel}/HIFI observations of high spectral resolution for FIR/submm molecular lines in two red supergiant stars (RSGs), NML Cyg and Betelgeuse, and two yellow hypergiant stars (YHGs), IRC+10420 and AFGL\\,2343. Tables~\\ref{tablines1} and~\\ref{tablines2} summarize the observational parameters derived from our data, which are displayed in Figs.~\\ref{fig_nml_h2o} to~\\ref{fig_afgl} -- the full spectra are shown in the appendix.  As illustrated in Fig.~\\ref{fig_histo}, the RSGs exhibit in general more intense high-excitation lines, indicating that they harbour inner circumstellar layers with high temperatures close to or even higher than 1000 K (Sections~\\ref{nmlcyg} and~\\ref{bet}). In contrast, YHGs do not contain this central warm material, and are instead surrounded by detached envelopes with low densities at small distances from the star (Sections~\\ref{irc} and~\\ref{afgl}). As such, our results show that the mass-loss rates in YHGs are presently weak. We also observe that the spectral line profiles in RSGs tend to become narrower with increasing excitation temperature. We show that this phenomenon indicates that the high excitation lines in RSGs originate from gas layers at small radii that have not yet reached the circumstellar terminal velocity. This trend is not observed in YHGs, in line with the picture of a hollow envelope closer to the stellar photosphere. These YHGs are thus less likely than RSGs to contain gas still in the acceleration phase.  In both YHGs, we detected intense emission features that are conspicuous in the profiles. That of IRC+10420 is relatively blue-shifted, at v$_{\\rm LSR}$ $\\sim$ 65 \\kms, and is more noticeable in certain molecules, such as H$_2$O and NH$_3$, although it does not seem to require particular excitation conditions. It was detected and mapped in low-$J$ CO Lines by CC07, who concluded that it comes from a condensation in the outer envelope. We also found a peculiar emission excess in AFGL\\,2343 (Section~\\ref{afgl}), at an extreme positive velocity of v$_{\\rm LSR}$ $\\sim$ 125 \\kms. There, the feature is clearly associated with the excitation of the lines. We speculate that this shell could result from a shock interaction between the circumstellar envelope and nearby interestellar gas.  We stress the particularly strong emission of water lines in NML Cyg and IRC+10420, which are also the only two sources featuring NH$_3$ emission. On the other hand, OH is ubiquitously observed. It is unclear where this molecule predominantly arises and how it forms in the envelopes. Its emission does not seem to be correlated with that of \\water, although OH formation in evolved stars is thought to result from H$_2$O photodissociation; for instance, the observed OH profiles are usually closer to those of mid-excitation CO lines and the OH line intensity is high in AFGL\\,2343, which shows particularly weak water lines.  Finally, we have proposed a preliminary model to fit the \\doce\\ and \\trece\\ lines detected in IRC+10420. The model is based on that developed by CC07 to explain their mm-wave interferometric mapping. We found that the original description of the shell around IRC+10420 by CC07 is able to reproduce all transitions by simply decreasing the temperature of the inner layers by about 30\\% (Fig.~\\ref{fig_irc_model_profile}). In addition, the emission from the high-$J$ lines is found to originate only from a detached, hot shell close to the star, formed by stellar mass ejection at a high rate of $\\sim$ 3$\\times$10$^{-4}$ \\my. Our calculation indicates that no hotter components at small radii are needed to reproduce the molecular emission of IRC+10420, confirming that the heavy mass-loss ceased about 200 years ago. More detailed modelling of the CO lines observed in IRC+10420 and AFGL\\,2343 will be presented in a forthcoming paper, in particular to take into account the departure from spherical symmetry illustrated by high-resolution maps at various wavelengths, and corroborated by the distinct spectral features detected in the HIFI data."
    },
    "1208/1208.4625_arXiv.txt": {
        "abstract": "Observations show that radial metallicity gradients in disk galaxies are relatively shallow, if not flat, especially at large galactocentric distances and for galaxies in the high-redshift universe. Given that star formation and metal production are centrally concentrated, this requires a mechanism to redistribute metals. However, the nature of this mechanism is poorly understood, let alone quantified. To address this problem, we conduct magnetohydrodynamical simulations of a local shearing sheet of a thin, thermally unstable, gaseous disk driven by a background stellar spiral potential, including metals modeled as passive scalar fields.  Contrary to what a simple $\\alpha$ prescription for the gas disk would suggest, we find that turbulence driven by thermal instability is very efficient at mixing metals, regardless of the presence or absence of stellar spiral potentials or magnetic fields.  The timescale for homogenizing randomly distributed metals is comparable to or less than the local orbital time in the disk. This implies that turbulent mixing of metals is a significant process in the history of chemical evolution of disk galaxies. ",
        "introduction": "The spatial distribution of metals in disk galaxies is a crucial clue for understanding how galaxies formed and evolved over cosmic time. The past few decades have produced a wealth of observations of this property, including in our own Milky Way \\citep[e.g.,][]{HK10,BR11,LL11,CR12,YCF12}, in nearby galaxies \\citep[e.g.,][]{VE92,CC00,PVC04,KC11}, and in the high-redshift universe \\citep[e.g.,][]{CM10,JE12,QC12}. In the local universe, a variety of radial metallicity gradients in disk galaxies are seen, but they are generally on the order of $-0.03$~dex~kpc$^{-1}$, with the negative sign indicating decreasing metallicity at larger galactocentric radii. However, surprisingly, these gradients seem to disappear in the outer parts of galactic disks, where there is little star formation; moreover, the metal content is greater than would be expected given the amount of star formation that has taken place at these radii \\citep[e.g.,][]{BR09,BKR12,WP11}.  High-redshift galaxies, in comparison, are far less regular. Their metallicity gradients range from negative ones significantly steeper than those found locally, to completely flat or even positive. Any successful theory of galactic evolution must be able to reproduce these observations.  Several processes play an important role in regulating the spatial variation of metals in disk galaxies, and these have been demonstrated by either chemical evolution models or hydrodynamical simulations.  The enrichment of metals in the interstellar medium (ISM) is dominated by star formation and subsequent stellar mass loss, and thus depends on the star formation law \\citep[e.g.,][]{PE91}.  Metals are diluted by infalling gas from outside galaxies \\citep[e.g.,][]{TL78,cC80,MF89,CMG97,PB00,CMR01}.  Radial inflow of the gas within the disk of a galaxy redistributes metals \\citep{MV81,LF85,PT89, GK92,PC00,SM11,BS12}; galaxy interactions are especially effective in inducing large-scale inflow and flattening metallicity gradients (\\citealt{RKC10,PMT11}; Torrey et al., in preparation). Turbulence associated with the viscous evolution of gas disks can also redistribute metals \\citep{cC89,SY89,SY90,TY95,TM98}. Beyond these gas-dynamical processes, radial migration of stars can alter stellar metallicity distributions independently of processes affecting the gas phase \\citep{SB02,RD08a,RD08b,SB09}.  However, the strength and relative importance of these processes remains very poorly understood. Semi-analytic chemical evolution models generally parameterize each process and then tune the parameters in an attempt to provide an acceptable match to observations. However, the large numbers of parameters involved means that even a good fit to the data may not be unique, and the need for fine-tuning means these models have limited predictive power. Moreover, the parameterizations used in the models may not be accurate. For example, turbulent mixing is usually treated by adopting an $\\alpha$ prescription for the turbulent transport of angular momentum \\citep{SS73}, and assuming that the transport coefficient for metals is the same. Under these assumptions turbulent mixing is unimportant unless $\\alpha$ is so large that the viscous diffusion time becomes comparable to the gas depletion time. However, neither the assumption that turbulent transport of metals can be approximated with an $\\alpha$ prescription, nor that $\\alpha$ for this process is the same as that for the angular momentum, are physically well-motivated. Numerical simulations of metal transport that evolve galaxies over cosmological times are unfortunately little better, because their limited resolution means that they must also adopt parameterized treatments of unresolved processes. For example, most smoothed particle hydrodynamics (SPH) simulations allow no chemical mixing at all between SPH particles \\citep{WVC08}, or at best treat mixing approximately using a parameterized subgrid recipe \\citep{SWS10}. Eulerian simulations, in contrast, dramatically over-mix when their resolution is low.  The approach we take in this paper is quite different, and complementary to chemical models and large-scale cosmological simulations. We isolate a single process: turbulent mixing within a galactic disk. Our goal is to provide a first-principles calculation of this process, which can in turn provide a physically-motivated, parameter-free prescription that can be used in chemical evolution models or lower resolution simulations. In order to achieve this goal, we simulate turbulent mixing in a portion of a galaxy at very high resolution, including physical processes that are too small-scale to be resolved in cosmological simulations, and we perform a resolution study to ensure that our results are converged. The previous work that most closely matches ours in philosophy and overall approach is that of \\citet{MF99} and \\citet{FML04}, who used high resolution simulations of isolated portions of galaxies to study mixing of supernova ejecta with the ISM as galactic winds are launched. Here we perform a similar calculation for turbulent mixing within disks.  There exist many sources that can drive turbulence in the ISM (see \\citealt{ES04} and \\citealt{SE04} and references therein).  In earlier work, \\citet{AM02} studied the properties of turbulent mixing driven by supernova explosions.  Here, we instead focus on turbulence driven by thermal instability \\citep{gF65, FGH69}. Our motivation is two-fold. First, the flat metallicity gradients seen in outer disks presumably call for some sort of mixing process to operate at large galactic radii, where star formation is limited and thus supernova explosions are extremely rare. Second, even in places where supernovae do occur, thermal instability is also present and will drive turbulence; indeed, thermal instability is essentially inevitable anywhere the ISM is dominated by atomic hydrogen, which is the case for most galaxies over the great majority of cosmic time. Supernovae will only enhance the turbulence compared to what we find, and thus our results should be viewed as a minimum estimate of the turbulent mixing rate.  In the sections that follow, we describe our numerical method and present the results of the simulations. We then quantify the results in a form appropriate for use in chemical evolution models and discuss their implications. Finally, we summarize the results and conclude. ",
        "conclusions": "In this work, we simulate a local patch of a vertically thin disk galaxy and study the transport of metals  with a variety of physical conditions.  Specifically, we investigate the ability of thermal instability, spiral shocks, and/or magnetic fields to homogenize metals.  We find that turbulence driven by thermal instability is especially effective in mixing the metals, regardless of the presence or absence of spiral shocks and magnetic fields.  We observe two different modes of turbulent mixing in our thermally unstable disks.  The first mode is for the turbulent gas to stir large-scale variations of metals into a random distribution.  The timescale for this mode is short compared to the local orbital time in the galaxy, and this mode may contribute to obliterate the chemical inhomogeneities introduced by star forming activities along spiral arms.  The second mode is for randomly-distributed metals to be continually homogenized over time by the turbulence.  We find the timescale for this process is relatively insensitive to wavelength and is on the order of half the orbital timescale.  This mode of turbulent mixing, therefore, should be of  significance in reducing the metallicity gradient in a disk galaxy.  We find that turbulent mixing of metals driven by thermal instability is more efficient than what a simple \\citet{SS73} $\\alpha$ prescription of viscosity for the gas would suggest.  The convective motion of the turbulent gas can in fact transport metals over larger distances, especially for kpc-scale variations.  The dynamics is perhaps more complicated than ordinary diffusive transport with a constant coefficient.  In an attempt to capture its qualitative behavior, however, we have devised a toy prescription in terms of a wavelength-dependent diffusion coefficient and measured its numerical values for our model galactic disk.  In principle, this prescription could be adopted as a sub-grid physical process in semi-analytic chemical evolution models as well as cosmological simulations.  Doing so should help us further constrain the dynamical history of disk galaxies."
    },
    "1208/1208.4908_arXiv.txt": {
        "abstract": "The observations of the \\textit{Kepler} space telescope revealed that fundamental-mode RR\\,Lyrae stars may show various radial overtones. The presence of multiple radial modes may allow us to conduct nonlinear asteroseismology: comparison of mode amplitudes and frequency shifts between observations and models. Here we report the detection of three radial modes in the star RR\\,Lyr, the eponym of the class, using the \\textit{Kepler} short cadence data: besides the fundamental mode, both the first and the ninth overtones can be derived from the data set. RR\\,Lyrae shows period doubling, but switches occasionally to a state where a pattern of six pulsation cycles repeats instead of two. We found hydrodynamic models that show the same three modes and the period-six state, allowing for comparison with the observations. ",
        "introduction": "The precise and continuous observations of the \\textit{Kepler} space telescope \\citep{borucki} revealed a wealth of new features in RR\\,Lyrae stars. Among the first discoveries was the detection of period doubling in three stars, including RR\\,Lyrae itself \\citep{kolenberg10, pd}. Additional frequencies in the Fourier transform, significant frequency components beside the usual RR\\,Lyrae pattern (main period, its harmonics, modulation sidelobes and the Blazhko-modulation frequency) were also detected in several stars \\citep{benko10, gug12}. Out of these new features, the period doubling phenomenon, the alternation of a higher and a lower amplitude pulsation cycle, has been successfully modeled. Our hydrodynamic calculations revealed that a high-order, 9:2 resonance can occur between the fundamental mode and the 9th radial overtone strange mode that leads to the period-doubling bifurcation \\citep{kmsz11}. These results were also confirmed by \\citet{smolec11}.  However, the nature of the additional frequencies is not fully understood yet, except for the ninth overtone which is related to the half-integer peaks and causes the period doubling. Some of the remaining peaks fall into the ranges where the first or second radial overtones are expected, but several stars display even more peaks, possibly indicating non-radial modes \\citep{benko10, gug12}. Additional modes were also detected in stars observed by the CoRoT space telescope \\citep{chadid10, poretti10, gug11}. The stars showing first or second overtone signals are unlike the usual double-mode RRd pulsators. The additional components fall into the mmag range so the amplitude ratios are extreme.  Furthermore, almost all stars that show additional modes are Blazhko-variables. Only two stars have been found in the \\textit{Kepler} sample that show a small second overtone with only traces of modulation down to the mmag level \\citep{benko10, nemec11}. On the other hand, about half of the modulated stars in the Kepler sample show additional peaks in the Fourier spectrum, often together with signs of period doubling. Period doubling itself is not a regular process: it is modulated and sometimes shows signatures of additional bifurcations. The variations do not necessarily follow the Blazhko-modulation either \\citep{pd}. Therefore it is necessary to investigate the origins of those variations.  \\begin{figure*} \\includegraphics[scale=.85]{f1.eps} \\caption{Short cadence Q6 data of RR\\,Lyrae. Light grey points are the brighter end of the light curve. Black and dark grey (blue and orange in the color version) lines connect the maxima of even and odd pulsation cycles respectively. In the case of the period-six visualization (shifted by 0.15 mag), every sixth maxima are connected, such as black/blue lines connect the $6k$, $6k+2$ and $6k+4$ points (circles), while grey/orange lines connect the $6k+1$, $6k+3$ and $6k+5$ points (dots) where $k=0,1,2$\\dots. The period-six-like sections are indicated with the grey boxes: we identified two occurrences between 55385.3-55406.3 and 55436.9-55448.3. Dashed lines mark the P6 set which we analyzed separately (see Figure \\ref{fou}). The insert shows the calculated luminosity of a model solution close to the 3:4 resonance. Note the similarity of the crossings between the individual blue and orange lines or branches. Model parameters ($\\alpha_\\nu$ is the eddy viscosity parameter) are: $M=0.64\\, M_\\odot$, $L=45\\, L_\\odot$, $T_{eff}= 6500\\, K$, $\\alpha_\\nu=0.0175$. The insert covers the same time span as the P6 section.  } \\label{q5-q6} \\end{figure*} ",
        "conclusions": "The detection of so many interesting features in a star of long observational history remarkably illustrates the advantages of space-based photometry. A decade ago, \\citet{smith03} noted that the residuals after fitting the primary components of the spectrum (the triplet at the fundamental mode, their harmonics and the modulation frequency) exceeded the expected observational error. The first observations of \\textit{Kepler} revealed the period-doubling phenomenon \\citep{kolenberg10, pd} and now the first overtone was also detected in the short cadence data. The results and implications of the analysis can be summarized as follows:  \\begin{itemize} \\item Three different radial modes were detected in RR\\,Lyr, the eponym of its class: the fundamental mode and the first and ninth overtone. However, the ninth overtone is locked in a resonance with the fundamental mode and can be detected only through the presence of period doubling \\citep{kmsz11}, causing the observed variations to resemble double-mode behavior. \\item The first overtone has a very small amplitude (few mmags), in contrast with the fundamental mode or the classical double-mode RRd stars, explaining the non-detection from the ground. \\item A state resembling period-six bifurcation was detected in the light variations of RR\\,Lyr. This behavior is most likely caused when the system approaches the vicinity of a 3:4 resonance between the fundamental mode and the first overtone. \\item We have been able to reproduce the three-mode state in nonlinear hydrodynamic models. Although we have not been able to model the 3:4 resonance or the higher than 0.750 period ratio of RR\\,Lyr, the amplitude ratios can be reproduced and higher-order resonances have been found as well. Furthermore, near-resonance models below the 0.750 value have very similar characteristics, raising the hopes for applying nonlinear asteroseismology to Blazhko RR\\,Lyrae stars. \\item Even the inclusion of three different modes cannot explain all the variations we see in the period doubling, indicating that the pulsation of RR\\,Lyrae stars may be even more complex. \\end{itemize}  Our radial nonlinear hydrodynamic calculations demonstrated that the period-doubled fundamental mode may lose its stability with respect to the first radial overtone. Then it is straightforward to expect similar nonlinear destabilization with respect to low order nonradial modes as well. This process might provide the background for the fact that additional pulsation frequencies occur more frequently in stars where period doubling is also detected. First overtone signals were identified in three other stars in the \\textit{Kepler} sample by \\citet{benko10}, and the case of RR\\,Lyr suggests that more may be discovered. The period ratios of peaks that can be plausibly associated with the first overtone are for V354 Lyr (KIC 6183128): $f_1/f_0=0.729$, and for V360 Lyr (KIC 9697825): $f_1/f_0=0.721$. For V445 Lyr (KIC 6186029), \\citet{benko10} noted a very complex pattern of peaks while the detailed analysis of \\citet{gug12} concluded that among the additional frequencies, two can be found at ratios $f_N/f_0=0.703$ and $f_0/f_1=0.730$. Even more stars show the signature of the second overtone for which \\citet{molnar12_f2} already proposed the alternate idea of a three-mode resonance, where the equation $3f_0+f_2=f_9$ may describe the relation between the three frequencies.  The results indicate that interactions between radial and nonradial modes may indeed play a more crucial role in the pulsation of stars. \\citet{bk11} showed that the resonance responsible for period doubling can cause Blazhko-like modulation in amplitude equations. Nonlinear asteroseismology promises to understand these new findings, and might even help to unlock the mystery of the Blazhko effect itself."
    },
    "1208/1208.5022_arXiv.txt": {
        "abstract": "Already more that 40 years ago, it has been suggested that because of the enormous mass densities in the cores of neutron stars, the hadrons in the centers of neutron stars may undergo a phase transition to deconfined quark matter. In this picture, neutron stars could contain cores made of pure (up, down, strange) quark matter which are surrounded by a mixed phase of quarks and hadrons. More than that, because of the competition between the Coulomb and the surface energies associated with the positively charged regions of nuclear matter and negatively charged regions of quark matter, the mixed phase may develop geometrical structures, similarly to what is expected of the sub-nuclear liquid-gas phase transition.  In this paper we restrict ourselves to considering the formation of rare phase blobs in the mixed quark-hadron phase. The influence of rare phase blobs on the thermal and transport properties of neutron star matter is investigated. The total specific heat, $c_V$, thermal conductivity, $\\kappa$, and electron-blob Bremsstrahlung neutrino emissivities, $\\epsilon_{\\nu,\\text{BR}}$, of quark-hybrid matter are computed and the results are compared with the associated thermal and transport properties of standard neutron star matter. Our results show that the contribution of rare phase blobs to the specific heat is negligibly small. This is different for the neutrino emissivity from electron-blob Bremsstrahlung scattering, which turns out to be of the same order of magnitude as the total contributions from other Bremsstrahlung processes for temperatures below about $10^8$~K. ",
        "introduction": "\\label{sec:intro}  Already many decades ago, it has been suggested that, because of the extreme densities reached in the cores of neutron stars, neutrons and protons may transform to quark matter in the cores of such objects \\cite{ivanenko65:a,fritzsch73:a,baym76:a,keister76:a,chap77:a+b,fech78:a}. Quark matter could thus exist as a permanent component of matter in the ultra-dense centers of neutron stars (see \\cite{glendenning01:a,glendenning00:book,weber99:book,weber05:a,maruyama07:a} and references therein). If the dense interior of a neutron star is indeed converted to quark matter, it must be three-flavor quark matter since it has lower energy than two-flavor quark matter. And just as for the hyperon content of neutron stars, strangeness is not conserved on macroscopic time scales, which allows neutron stars to convert confined hadronic matter to three-flavor quark matter until equilibrium brings this process to a halt.  As first realized by Glendenning \\cite{glendenning92:a}, the presence of quark matter enables the hadronic regions of the mixed phase to arrange to be more isospin symmetric than in the pure phase by transferring charge to the quark phase in equilibrium with it. The symmetry energy will be lowered thereby at only a small cost in rearranging the quark Fermi surfaces. The electrons play only a minor role when neutrality can be achieved among the baryon-charge carrying particles. The stellar implication of this charge rearrangement is that the mixed phase region of the star will have positively charged regions of nuclear matter and negatively charged regions of quark matter.  Because of the competition between the Coulomb and the surface energies associated with the positively charged regions of nuclear matter and negatively charged regions of quark matter, the mixed phase may develop geometrical structures (see Fig.\\ \\ref{fig:qh_structures}), similarly as it is expected of the subnuclear liquid-gas phase transition \\cite{ravenhall83:a,ravenhall83:b,williams85:a}. This competition establishes the shapes, sizes, and spacings of the rare phase in the background of the other in order to minimize the lattice energy \\cite{glendenning92:a,maruyama07:a,glendenning01:a,glen95:a}.  The change in energy accompanied by developing such geometrical \\begin{figure} \\includegraphics[scale=0.2]{f1.eps} \\caption{Schematic illustration of possible geometrical structures in the quark-hadron mixed phase of neutron stars. The structures may form because of the competition between the Coulomb and the surface energies associated with the positively charged regions of nuclear matter and negatively charged regions of quark matter.} \\label{fig:qh_structures} \\end{figure} structures is likely to be very small in comparison with the volume energy \\cite{glendenning92:a,glen01:b,heiselberg92:a,heiselberg95:crete} and, thus, may not much affect the global properties of a neutron star. However, the geometrical structure of the mixed phase may be very important for irregularities (glitches) in the timing structure of pulsar spin-down as well as for the thermal and transport properties of neutron stars \\cite{glendenning92:a,glen01:b,glendenning00:book}.  To calculate the neutrino-pair bremsstrahlung rates and thermal properties, we follow the method described in \\cite{kaminker99:a} and \\cite{potekhin99:a}, which is commonly used for the calculation of the neutrino emissivity and thermal conductivity in the crusts of neutron stars.  These authors considered contributions from electron-phonon scattering and Bragg diffraction (the static-lattice contribution). Furthermore, multi-phonon processes and electron band structure effects are incorporated to obtain more realistic scattering rates and a better connection between the solid and the liquid gas phase.  Instead of adopting the analytic fits provided in \\cite{kaminker99:a} and \\cite{potekhin99:a}, here we re-calculate the scattering rates from phonon sums using the method of \\cite{mochkovitch79:a}. There are two main reasons for this. The first being that, for the crust, the total ion charge is balanced by the total electron charge. This will be different for the mixed quark-hadron phase in the core of a neutron star, since electric charge neutrality is established between the electric charges of the rare phase, the dominant phase, and the leptons which are present in both the rare and the dominant phase. The simple relation $n_e=Z n_i$ between electron density and ion density, used to derive the crustal fit formula in \\cite{kaminker99:a,potekhin99:a}, can therefore not be used to study the quark-hadron Coulomb lattice structure in the core of a neutron star. The second reason concerns the electric charge numbers themselves. For mixed phase blobs, they can easily exceed $Z\\sim10^3$, as will be shown in \\S~\\ref{sec:mixedph}. Charge numbers that high are obviously not reached in the crustal regimes of neutron stars \\cite{kaminker99:a}, where there is usually no need to consider atomic nuclei with charges much larger than $Z>56$.  The paper is organized as follows. In Section~\\ref{sec2}, we briefly discuss the modeling of the mixed quark-hadron phase in the cores of neutron stars and the equations of state of confined hadronic and quark matter used in this work. In Section~\\ref{sec:calc}, we summarize the formalism for calculating the neutrino-pair Bremsstrahlung emissivity and the thermal conductivity of rare phase blobs immersed in hadronic matter. The results are presented in Section~\\ref{sec:result}. ",
        "conclusions": "\\label{sec:conclusion}  Because of the competition between the Coulomb and the surface energies associated with the positively charged regions of nuclear matter and negatively charged regions of quark matter, the mixed phase may develop geometrical structures (e.g., blobs, rods, slabs), similarly to what is expected of the sub-nuclear liquid-gas phase transition.  In this paper we explore the consequences of a Coulomb lattice made of rare phase blobs for the thermal and transport properties of neutron stars. The total specific heat, $c_V$, thermal conductivity, $\\kappa$, and electron-blob Bremsstrahlung neutrino emissivities, $\\epsilon_{\\nu,\\text{BR}}$, are calculated and compared with those of standard neutron star matter.  To carry out this project, we have adopted, and expanded on, methods of earlier works on the transport properties of neutron stars \\cite{glendenning89:a,glendenning01:a}.  The sizes of, and spacings between, rare phase blobs are calculated using the Wigner-Seitz approximation \\cite{glendenning01:a}.  The equations of state used in this study are computed for a standard non-linear nuclear Lagrangian, and the associated equations of motion for the baryon and meson fields are solved in the relativistic mean-field approximation. Quark matter has been modeled in the framework of the MIT bag model.  Four different parameter sets (HV1, HV2, G300I, G300II) have been used to model the composition of neutron star matter containing a mixed phase of quarks and hadrons (quark-hybrid matter).   The results discussed in Section~\\ref{sec:result} show that the contribution of rare phase blobs in the mixed phase to the specific heat is negligible compared to the specific heat of a quark-hadron gas. This is very different for the transport properties.  For low temperature $T\\lesssim 10^8$~K the neutrino emissivity from electron-blob Bremsstrahlung scattering is at least as important as the total contribution from other Bremsstrahlung processes (such as nucleon-nucleon and quark-quark Bremsstrahlung) and modified nucleon and quark Urca processes (see Figs.\\ \\ref{fQHV} and \\ref{fQG300}).  It is also worth noting that the scattering of degenerate electrons off rare phase blobs in the mixed phase region lowers the thermal conductivity by several orders of magnitude compared to a quark-hadron phase without geometric patterns (see Figs.\\ \\ref{fkappa1} through \\ref{fkappa4}).  This may lead to significant changes in the thermal evolution of the neutron stars containing solid quark-hadron cores, which will be part of a future study. Another very interesting issue concerns the impact of more complex geometrical structures (rods and slabs) on the thermal conductivity and on neutrino transport. The presence of such structures may reduce the neutrino emissivities because of changes in the dimension of the reciprocal lattice and the Debye-Waller factor \\cite{kaminker99:a}.  In summary, our study has shown that the presence of rare phase blobs in dense neutron star matter may have very important consequences for the total neutrino emissivity and thermal conductivity of such matter. The implications of this for the thermal evolution of neutron stars need to be explored in future studies. To accomplish this we intend on performing two dimensional cooling simulations, in which rotation and a dynamic composition might be accounted for \\cite{Negreiros:2011a,Negreiros:2012a}. In this connection we refer to the recent study of Noda {\\it et al.}\\ \\cite{noda11}, who suggested that the rapid cooling of the neutron star in Cassiopeia A can be explained by the existence of a mixed quark-hadron phase in the center of this object."
    },
    "1208/1208.5487_arXiv.txt": {
        "abstract": "The inspiral and merger of a binary neutron star (NSNS) can lead to the formation of a  hypermassive neutron star (HMNS). As the HMNS loses thermal pressure due to neutrino cooling and/or centrifugal support due to gravitational wave (GW) emission, and/or magnetic breaking of differential rotation it will collapse to a black hole. To assess the importance of shock-induced thermal pressure and cooling, we adopt an idealized equation of state and perform NSNS simulations in full GR through late inspiral, merger, and HMNS formation, accounting for cooling. We show that thermal pressure contributes significantly to the support of the  HMNS against collapse and that thermal cooling accelerates its ``delayed'' collapse. Our simulations demonstrate explicitly that cooling can induce the catastrophic collapse of a {\\it hot} hypermassive neutron star formed following the merger of binary neutron stars. Thus, cooling physics is important to include in NSNS merger calculations to accurately determine the lifetime of the HMNS remnant and to extract information about the NS equation of state, cooling mechanisms, bar instabilities and B-fields from the GWs emitted during the transient phase prior to BH formation. ",
        "introduction": "The inspiral and merger of compact binaries has attracted considerable attention in recent years for two main reasons. First, such systems emit a large flux of gravitational waves (GWs), making them among the most promising sources for GWs detectable by ground-based laser interferometers such as LIGO \\cite{LIGO1,LIGO2}, VIRGO \\cite{VIRGO1,VIRGO2}, GEO \\cite{GEO}, and KAGRA \\cite{KAGRA}, as well as by proposed space-based interferometers such as eLISA/NGO \\cite{NGO} and DECIGO \\cite{DECIGO}. Second, black hole -- neutron star (BHNS) and neutron star -- neutron star (NSNS) mergers are candidates for the central engines that power the observed short-hard gamma ray bursts (sGRBs).  Extracting physical information about these binaries from their GWs and their accompanying electromagnetic signals may reveal critical details about the equation of state of neutron star matter and may unveil the nature of the sGRB phenomenon. However, interpreting the data requires careful modeling of these systems in full general relativity (see \\cite{BSBook} for a comprehensive review and references). Most effort in general relativity to date has focused on modeling black hole--black hole (BHBH) binaries (see also \\cite{Hinder:2010vn}), and neutron star--neutron star (NSNS) binaries (see also \\cite{BNSlr}), with some recent work on black hole--neutron star binaries (see also \\cite{st11}), and white dwarf--neutron star binaries \\cite{Paschalidis:2010dh,PLES2011,Paschalidis:2009zz}.   NSNSs are known to exist, which makes NSNS systems particularly attractive to study. Theoretical calculations show that NSNS mergers can lead to the formation of a hypermassive neutron star. A HMNS \\cite{HMNS} is a differentially rotating NS whose mass exceeds the maximum mass of a uniformly rotating star \\cite{CST1994a,CST1994b}. The latter is about 20\\% larger than the maximum mass of a nonrotating (spherical) equilibrium star (the TOV limit) \\cite{HMNS}. Typically a HMNS forms following the merger of a NSNS, when the system's total mass is smaller than some threshold mass $M_{\\rm th}$. According to \\cite{ST} this threshold mass is $M_{\\rm th} \\approx 1.3-1.35 M_{\\rm sph}$, where $M_{\\rm sph}$ is the TOV limit for the same EOS.  A HMNS is a transient, quasiequilibrium configuration. It will eventually undergo ``delayed collapse'' on a secular (dissipative) time scale, which may power a sGRB. There are two distinct routes by which this collapse might be triggered: \\begin{enumerate}  \\item If the HMNS is primarily centrifugally supported, redistribution of angular momentum by viscosity or magnetic fields \\cite{DLSS2004,dlsss06a}, and/or loss of angular momentum by GW emission \\cite{SHTA2006} destroys the support provided, leading to catastrophic collapse.  \\item If the HMNS is primarily supported by thermal pressure generated by shocks during merger, delayed collapse may be triggered by the loss via neutrino cooling of thermal energy~\\footnote{Note that neutrinos too carry away angular momentum from the system, but according to \\cite{Baumgarte:1998sn} neutrino emission is very inefficient in decreasing the angular momentum of a HMNS.}. \\end{enumerate}  While catastrophic collapse of a {\\it cold} HMNS via viscosity or magnetic fields has been demonstrated using fully general relativistic calculations \\cite{DLSS2004,dlsss06a}, there are no fully general relativistic calculations to date that demonstrate explicitly that cooling can induce collapse of a {\\it hot} HMNS produced following the merger of binary neutron stars.  HMNSs formed in NSNS mergers will always be {\\it hot} due to shock heating. A priori it is not clear which mechanism is most important for holding up a  HMNS against collapse: {\\it centrifugal forces} or {\\it thermal pressure}. The answer to this question is still open and may depend on the nature of the companions (e.g. masses, EOS etc.).  Recent simulations of binary NS mergers that form hypermassive NSs seem to point in different directions. For example, in \\cite{2008PhRvD..78h4033B,Rezzolla:2010fd} an equal-mass NSNS is evolved assuming a $\\Gamma=2$ equation of state (EOS). It is shown that angular momentum carried away by gravitational waves alone can induce the collapse. Reference \\cite{2011PhRvL.107e1102S} also evolves an equal-mass NSNS, but with a more realistic, finite temperature, nuclear EOS. They find that the deviation of their HMNSs from axisymmetry is so small that GW emission is significantly reduced. The authors argue that shock heating is sufficiently important that their  HMNSs are supported by the excess thermal pressure.  Determining which mechanism controls the lifetime of the remnant is important because it determines the time interval between the NSNS merger and the delayed collapse -- a time interval that can in principle be measured by Advanced LIGO/VIRGO. It is the time interval between the end of the gravitational wave signal due to the inspiral and the beginning of the burst signal due to the delayed collapse. If differential rotation support is most important, then the time interval is governed by, e.g. the Alfv\\'{e}n time scale, assuming magnetic braking of differential rotation is most important, or the GW time scale, in the case of a rapidly spinning remnant that develops a bar. By contrast, if thermal pressure is dominant, then the time scale is governed by thermal cooling. Therefore, knowing the mechanism driving collapse may place constraints on seed magnetic field magnitudes, or the existence of bar modes, or the relevant cooling mechanisms. It could even place constraints on the temperature of matter, as well as the nuclear EOS.  To disentangle the effects of thermal support from those of rotational support, previous studies compared results from NSNS simulations that suppress shocks (by enforcing a strictly cold EOS) to those that allow shocks. If the HMNS remnant lives longer with shocks than without, then it is tempting to infer that thermal pressure due to shock heating is chiefly responsible for supporting the remnant. However it is not possible to draw such a firm conclusion because shocks, which act on a hydrodynamical time scale, not only heat the gas, thereby increasing the total pressure support, but also affect the matter and angular momentum profiles. Different profiles can themselves increase the lifetime of a HMNS.  The goal of this paper is to study the relative importance of thermal pressure in supporting  HMNSs from collapse and demonstrate that cooling can induce the catastrophic collapse of a  HMNS formed following the merger of binary neutron stars. We accomplish this by performing a limited set of NSNS simulations in full GR through late inspiral, merger, (hot) HMNS formation, and collapse. We account for cooling in the HMNS remnant via a covariant cooling scheme we developed in \\cite{PLES2011}. We then compare this HMNS evolution to a control simulation, in which the cooling mechanism is disabled.  Our simulations model the initial NSNS binary as equal-mass, irrotational, quasiequilibrium $n=1$ polytropes in a quasicircular orbit, corresponding to case \\mbox{$1.46$-$45$-$\\ast$} of \\cite{2008PhRvD..78h4033B}.  Following the NSNS merger, a quasiequilibrium HMNS forms. We then continue the evolution of the remnant with and without cooling, which we model via an effective local emissivity. For the runs with cooling we choose two cooling time scales. We find that, independent of the cooling time scale chosen, the HMNS collapses and forms a BH within a few cooling time scales.  Our simulations suggest that shock-induced thermal pressure is a significant source of support against gravitational collapse, even in the case of polytropic NSs and demonstrate explicitly that cooling can induce the catastrophic collapse of a  HMNS. Estimating the temperature of the remnant, we find that a realistic neutrino cooling time scale is of order a few $100$ms. Given that our estimated cooling time scale is comparable to the angular momentum redistribution/loss time scales due to either magnetic braking or GWs, our results suggest that accounting for cooling is a critical ingredient in predicting the lifetime of a  HMNS. Accordingly, cooling physics must be incorporated in models of binary NS simulations.  The paper is structured as follows. In Sec.~\\ref{sec:timescales} we review the time scales relevant to HMNSs formed in binary NSNS mergers. Sections~\\ref{sec:basic_eqns} and~\\ref{sec:numerical} summarize the initial data, basic evolution equations, numerical methods, and cooling formalism. The basic results are presented in Sec.~\\ref{sec:results} and summarized in Sec. \\ref{sec:summaryandfuturework}. Throughout this work, geometrized units are adopted, where $G = c = 1$, unless otherwise specified. ",
        "conclusions": "\\label{sec:summaryandfuturework}    A differentially rotating, quasiequilibrium HMNS is a transient configuration that can arise following the merger of NSNS binaries. The mass of a HMNS is larger than the maximum mass that can be supported by a cold EOS, even with maximal uniform rotation. A HMNS will eventually undergo ``delayed collapse'' on a secular (dissipative) time scale and may power a sGRB.  When HMNSs are born in NSNS mergers, they are rapidly differentially rotating and {\\it hot} due to shock heating. Therefore, HMNSs will collapse to a BH either on an angular momentum loss/magnetic braking time scale or on a cooling time scale. A priori it is not clear which of the two above mechanisms is most important for holding up an HMNS against collapse: {\\it centrifugal forces or thermal pressure}. The answer to this question is still open and may depend on the stellar model, EOS, and initial magnetic fields.  Determining which mechanism drives a  HMNS to collapse has observational consequences; the time scale of collapse will set the interval between the NSNS merger chirp signal and the delayed collapse burst signal, which may be measured by LIGO/VIRGO. Careful modeling of HMNS physics will thus place constraints on magnetic field magnitudes, the existence of bar modes, and/or the relevant cooling mechanisms. In addition, such observations could place constraints on the temperature of matter as well as the nuclear EOS.  To disentangle the effects of thermal support from those of rotational support, previous studies compared results from NSNS simulations that suppress shocks to those that allow shocks. If the HMNS remnant lives longer in the case with shocks than without, then it is tempting to infer that thermal pressure due to shock heating is solely responsible. However it is not possible to draw such a firm conclusion because shocks, which act on a hydrodynamical time scale, not only heat the gas, thereby increasing the total pressure support, but also affect the matter profile and redistribute angular momentum. Different matter and angular momentum profiles alone can increase the lifetime of a HMNS via an increase of both the GW time scale and the amount of differential rotational support.  To address this issue, we first performed long-term, high-resolution GRHD NSNS simulations through inspiral, merger, and HMNS formation, allowing for shocks. Following HMNS formation, we continue the evolution both with and without cooling. When cooling is turned off, the remnant collapses on the GW time scale. However, when cooling is turned on we find that the HMNS collapses and forms a BH within a few cooling time scales.  Our simulations demonstrate that shock-induced thermal pressure is a {\\it significant} source of support against gravitational collapse in the case of a stiff $\\Gamma$-law EOS -- a result consistent with simulations that employ a more realistic EOS~\\cite{2011PhRvL.107e1102S} -- and show explicitly that cooling can induce the catastrophic collapse of a HMNS. Estimating the temperature of the HMNS remnant, we find that a realistic neutrino cooling time scale is of order a few $100$ms. Given that the estimated cooling and angular momentum loss/magnetic braking time scales can be comparable, cooling should be accounted for to accurately determine the lifetime of a HMNS. Therefore simulations that implement cooling will lead to earlier collapse than simulations that ignore it, otherwise the predicted GW and EM signatures from these delayed collapse events may be incorrect.  Therefore, to accurately determine the lifetime of HMNS remnants, neutrino cooling physics should be incorporated in NSNS simulations. In the future we plan to revisit the subject using a more realistic neutrino leakage scheme, such as that used in \\cite{Sekiguchi:2010fh,2011PhRvL.107e1102S}, in conjunction with more realistic treatment of the microphysics involved."
    },
    "1208/1208.4979.txt": {
        "abstract": "{This paper deals with the cosmic-ray penetration into molecular clouds and with the related gamma--ray emission. High energy cosmic rays interact with the dense gas and produce neutral pions which in turn decay into two gamma rays. This makes molecular clouds potential sources of gamma rays, especially if they are located in the vicinity of a powerful accelerator that injects cosmic rays in the interstellar medium. The amplitude and duration in time of the cosmic--ray overdensity around a given source depend on how quickly cosmic rays diffuse in the turbulent galactic magnetic field. For these reasons, gamma-ray observations of molecular clouds can be used both to locate the sources of cosmic rays and to constrain the properties of cosmic-ray diffusion in the Galaxy.} ",
        "introduction": "\\label{sec:intro}  Cosmic Rays (CRs) \\cite{berezinskii,gaisser,schlickeiser,longair} are charged and energetic particles that hit the Earth's atmosphere from above. The flux of CRs, once corrected for the effect of solar modulation, is constant in time and corresponds to a local energy density of $w_{CR} \\approx 1$~eV/cm$^3$. Remarkably, this is comparable to the energy density of both magnetic field and thermal gas in the interstellar medium. CRs are mainly protons, with approximately 10\\% of Helium, and 1\\% of both heavier nuclei and electrons. Their differential energy spectrum is a steep and featureless power law $N_{CR} \\propto E^{-s}$ with slope $s \\approx 2.7$, and $\\approx$~GeV particles are the main contributors to the total CR energy density. The slope of the spectrum slightly steepens  to $s \\approx 3$ at an energy of $\\approx 4 \\times 10^{15}$~eV and this spectral feature is called the CR knee. The CR spectrum continues up to energies of the order of $\\approx 10^{20}$~eV, but here we restrict ourselves to considering particles with energies below the knee, which can be confined by the interstellar magnetic field and thus are certainly of galactic origin. Finally, the arrival directions of CRs are extremely isotropic in the sky. The isotropy is of the order of $\\approx 10^{-3}$ for particle energies  above $\\approx$~1~TeV, where local (i.e. heliospheric) effects can be neglected, and depends very weakly on the particle energy for energies up to the knee. The high level of isotropy is due to the diffusion of CRs in the turbulent galactic magnetic field, which isotropizes the trajectories of particles and prevents a direct identification of CR sources based on the observed arrival direction of particles. This is the reason why indirect observational evidences, such as the detection of photons produced by CR interactions with the ambient medium, are needed in order to locate the sites of CR acceleration. {\\it To date, the sources of galactic CRs are still not firmly identified.}  A connection between CRs and supernovae was first proposed by Baade and Zwicky in 1934 \\cite{baadezwicky} and still remains the most popular explanation for the origin of galactic CRs. In its modern version (see \\cite{hillas} for a review), the supernova paradigm for the origin of cosmic rays mainly relies on a consideration based on  the energy required to maintain the observed flux of CRs against their escape from the Galaxy. The overabundance of Li, Be, and B in CRs with respect to the abundances measured in the solar system, where they are virtually absent\\footnote{Li, Be, and B are not synthesized in stars, and their standard abundance is very low, being mainly determined by primordial nucleosynthesis.}, can be explained as the result of spallation of heavier CR nuclei by interstellar gas. The amount of matter, or grammage, that CRs with an energy of $\\gtrsim$GeV need to traverse to produce the observed amount of Li, Be, and B is equal to $\\mu \\approx 5$~g/cm$^2$. This corresponds to a confinement time in the galactic disk of $t_{d} = \\mu/\\rho c \\approx 3 \\times 10^6$~yr, where $\\rho$ is the mean gas density in the disk ($\\approx$~1 particle per cubic centimeter) and $c$ is the speed of light. Assuming that the CR intensity is constant in both time and space within the galactic disk, which has a radius of $R_{mw} \\approx$~15~kpc and a thickness $h$ of a few hundred parsecs, one can estimate the CR luminosity of the Galaxy as $W_{CR} = [w_{CR} (\\pi R_{mw}^2) h]/t_d \\approx 10^{41}$~erg/s. This has to be compared with the total power from supernova explosions in the galaxy $P_{SN} = \\nu_{SN} E_{SN} \\approx 10^{42}$~erg/s, where $\\nu_{SN} \\approx 3$/century is the supernova rate in the Galaxy and $E_{SN} \\approx 10^{51}$~erg is the typical supernova explosion energy. It is evident form these figures that {\\it supernovae, or something related to them, may be the sources of CRs if $\\approx$~10\\% of their explosion energy is somehow converted into accelerated particles.}  A mechanism for the acceleration of particles that operates at supernova remnant (SNR) shocks was proposed in the late seventies, when it was realized that particles can be accelerated at shock waves via a first-order Fermi mechanism \\cite{krimsky,bell78,blandfordostriker}. A characteristic prediction of these models is a differential energy spectrum for the accelerated particles which is a power law with slope close to $\\approx E^{-2}$. %If CRs are treated as test particles (i.e. their pressure is negligible and does not affect the shock structure) the differential spectrum of accelerated particles is a power law in energy $E^{-\\alpha}$ with a slope that depends uniquely on the shock compression factor $r$ as $\\alpha = (r+2)/(r-1)$. For high Mach number shocks, as SNR shocks are, $r \\rightarrow 4$ and the spectrum of accelerated particles converges towards the universal shape $E^{-2}$. %Non-linear theories for particle acceleration at shocks have been developed, to consider both the effect of the pressure of CRs onto the shock structure \\cite{malkovdrury} and the amplification of the magnetic field due to CR streaming \\cite{lagagecesarsky,bell04}. After combining these two effects, the $E^{-2}$-universality is broken, and particle spectra can be, according to model details, both harder (with $\\alpha$ up to $\\approx$1.5, see e.g. \\cite{malkovdrury}) or softer (values of the order of $\\alpha \\approx 2...2.4$ or so are found in the literature, e.g. \\cite{alfvendrift})  than that. Power law spectra of relativistic particles have indeed been observed in SNRs, both through X-ray (see e.g. \\cite{koyama,reynolds}) and gamma-ray (e.g. \\cite{felixreview,jimreview}) observations and this is considered an evident manifestation of shock acceleration at expanding SNR shock waves. The X-ray emission is unambiguously interpreted as the synchrotron radiation from relativistic electron, while the gamma ray emission can be interpreted either as inverse Compton scattering of electrons or decay of neutral pions generated in hadronic interactions between CRs and ambient gas. As discussed in the following, {\\it the ambiguity between the hadronic or leptonic origin of the observed gamma-ray emission from SNRs is one of the main obstacles in proving (or disproving) the fact that SNRs are indeed the sources of CRs.}  Another issue that needs to be explained is the absence of features in the CR spectrum up to the energy of the CR knee, which is the mild steepening of the spectral slope observed at a particle energy of a few PeVs. The featureless of the spectrum up to that energy suggests that the sources that are responsible for the acceleration of the bulk of the CRs in the Galaxy ($\\approx$GeV particles) are probably able to accelerate particles all the way up to the knee. In fact, assuming that several classes of sources contribute significantly to the CR spectrum at different particle energies and create such a featureless power law spectrum, though not impossible, would require an {\\it ad hoc} fine tuning which seems quite unreasonable. In other words, {\\it if SNRs are the sources of galactic CRs, most likely they have to act as particle PeVatrons} (see e.g. \\cite{gabiciPeV}).  Finally, the question on the origin of galactic CRs cannot be considered answered until we understand the details of their propagation in the interstellar medium (for recent reviews see e.g. \\cite{andyreview,fiorenzareview}). Measurements of the grammage that CRs must traverse while propagating from the sources to the Earth can be performed at different CR particle energies. Such measurements clearly point toward an energy dependent  grammage, and thus an energy dependent confinement time of CRs in the Galaxy, with higher energy particles escaping faster, according to $t_{esc}(E) \\propto E^{-\\delta}$, with $\\delta \\approx 0.3...0.6$. By assuming that CRs of all energies travel, on average, a distance $h$ before leaving the Galaxy, the escape time can be converted into a spatial diffusion coefficient $D \\approx h^2/t_{esc} \\approx D_0 (E/10~{\\rm GeV})^{\\delta}$, with $D_0 \\approx 10^{28}...10^{29}$~cm$^2$/s \\cite{andyreview,fiorenzareview}. If CR sources inject in the Galaxy $Q_{CR}(E)$ particles with energy $E$ per unit time, with a power law spectrum $Q_{CR}(E) \\propto E^{-\\alpha}$, then the equilibrium spectrum of CRs in the Galaxy is: $N_{CR} \\propto Q_{CR}(E) \\times t_{esc} \\propto E^{-\\alpha-\\delta}$. The observed slope of the CR spectrum is $N_{CR} \\propto E^{-2.7}$, which gives: $\\alpha \\approx 2.1...2.4$. Thus, {\\it the slope of the injection spectrum of CRs in the Galaxy has to be close to, but definitely steeper than 2}.  At this point, another remark is needed. If the diffusion coefficient grows too fast with energy, i.e. if $\\delta$ is closer to $\\approx$~0.6 rather than to $\\approx 0.3$, CRs with energies close to the knee would escape the Galaxy too quickly to be isotropized by the galactic magnetic field. In this scenario one would expect a high level of anisotropy at high energies, in contrast with what is observed. Thus, best--bet reference values for the spectral slope at injection of CRs and for the slope of the diffusion coefficient are probably $\\alpha \\approx 2.4$ and $\\delta \\approx 0.3$, which are consistent with both  the chemical and isotopic abundances of CRs and their isotropy.  The reason why the SNR hypothesis for the origin of galactic CRs is the most trusted and investigated scenario (but see e.g. \\cite{bubbles} and \\cite{dar} for a different and radically different perspective, respectively) is the fact that within this framework, most of the observational requirements can be explained within a reasonable accuracy. As said above, the total energy budget is not an issue, provided that the efficiency of particle acceleration is of the order of $\\approx$10\\%. X-ray and gamma-ray observations clearly show that a mechanism capable of accelerating particles up to (at least) hundreds of TeVs operates at SNR shocks, and the characteristics of the observed radiation fit quite well with predictions of shock acceleration theory. Moreover, recent developments in our understanding of the CR-induced amplification of the magnetic field at shocks suggest that SNRs might be able to accelerate particles up to the energy of the knee \\cite{bell04} and inject them in the interstellar medium with a spectrum slightly steeper than $E^{-2}$ (see e.g. \\cite{brianperp} or \\cite{alfvendrift,driftcaprioli} for two ways to steepen the spectral slope above $\\alpha = 2$), as required to explain the observed spectrum of CRs. Also the CR chemical composition is reproduced with fair agreement with observations \\cite{pzs}, while for what concerns the CR anisotropy the agreement between predictions and data is consistent within a factor of a few \\cite{ptuskinanisotropy,pasqualeanisotropy}, if a weak dependence on energy of the diffusion coefficient is adopted, i.e. $\\delta \\approx 0.3$. However, in this latter case a comparison is less straightforward given that the level of anisotropy may be dominated by the exact location of the few nearest CR sources.  Despite all these very encouraging facts, it has to be kept in mind that we are still missing a conclusive and unambiguous proof of the fact that SNRs, as a class of objects, accelerate CRs and inject them in the interstellar medium at the rate required by observations. This review is an attempt to describe how gamma-ray observations, and in particular gamma-ray observations of molecular clouds, might finally lead to prove (or falsify) the SNR paradigm for the origin of CRs. ",
        "conclusions": "In this paper, it has been shown how gamma--ray observations of MCs can be used to identify the location of the sources of galactic CRs and to constrain the CR diffusion properties close to the sources. Despite encouraging results from both observations and theory, further work is needed to reach a conclusive evidence in favor (or against) the SNR paradigm for the origin of CRs. Hopefully, this will come with the advent of the next generation of gamma--ray instruments, such as CTA.  The main limitations of the approach presented here are probably connected with the oversimplified assumptions made to describe the way in which particles escape from SNR shocks, which is still not well understood (see \\cite{gabiciescape} and references therein for a discussion) and the way in which they diffuse in the interstellar medium.  Of great relevance is the fact that the presence of a MC interacting with the SNR shock can, one one side, amplify the gamma--ray emission from neutral pion decay \\cite{adv}, but also influence and modify the acceleration mechanism of particles at shocks \\cite{lukeneutrals,brianneutrals,crushedyas}, and even have important effects on their escape \\cite{malkovbreak}. All these aspects need to be further investigated.  Moreover, the assumption of isotropic and homogeneous diffusion of CRs close to their sources is certainly an excessive oversimplification. In fact, CR are expected to diffuse preferentially along, parallel to, the magnetic field lines, the perpendicular transport being determined mainly by the wandering of the magnetic field line CRs are attached to \\cite{parker,fabien}. In adition to that, the CR diffusion along the magnetic field lines is most likely a non--linear process, where CRs themselves generate the magnetic turbulence needed to confine them. Some preliminary work including the effects of anisotropic CR diffusion have been recently published \\cite{plesser,michael,malkovfelix,lara}, and this promises to become one of the most important developments in this field.  Finally, recent measurements of the CR ionization rate in a MC interacting with the shock of the SNR W51C  have been presented \\cite{cecilia}. These data reveal an enhancement in the CR ionization rate of about 2 orders of magnitude with respect to standard values. Such an enhancement might be interpreted as the result of the presence of CRs accelerated at the SNR shock. The SNR/MC system has been also detected in TeV gamma rays \\cite{julian}, and the emission is most likely hadronic. Thus, in this particular system CRs can be studied from $\\approx$~MeV energies (the most relevant for the ionization of the gas) up to multi TeV energies. In the future, studies of this kind will shed light on the acceleration of CRs at shocks and on their escape over an unprecedented energy range.    \\begin{acknowledgement} I would like to thank the organizers of the Sant Cougat Forum for Astrophysics, Diego Torres and Olaf Reimer, for their invitation. I also acknowledge support from the EU [FP7--grant agr. n$^{\\circ}$256464] and from ANR [JCJC Programme]. \\end{acknowledgement} %"
    },
    "1208/1208.4953_arXiv.txt": {
        "abstract": "{The Magellanic Cloud System (MCS) interacts via tidal and drag forces with the Milky Way galaxy. } {Using the Parkes Galactic All--Sky Survey (GASS) of atomic hydrogen we explore the role of drag on the evolution of the so-called Leading Arm (LA). } {We present a new image recognition algorithm that allows us to differentiate features within a 3-D data cube (longitude, latitude, radial velocity) and to parameterize individual coherent structures. We compiled an \\hi object catalog of LA objects within an area of 70\\degr x 85\\degr (1.6 sr) of the LA region. This catalog\\thanks{Table 2 is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bonn/qcat?J/A+A/} comprises information of location, column density, line width, shape and asymmetries of the individual LA objects above the 4--$\\sigma$ threshold of $\\Delta T_b \\simeq 200\\,{\\rm mK}$. } {We present evidence of a fourth arm segment (LA4). For all LA objects we find an inverse correlation of velocities $v_{\\rm GSR}$ in Galactic Standard of Rest frame with Magellanic longitude. High--mass objects tend to have higher radial velocities than low--mass ones. About 1/4 of all LA objects can be characterized as head-tail (HT) structures. } {Using image recognition with objective criteria, it is feasible to isolate most of  LA emission from the diffuse Milky Way \\hi gas. Some blended gas components (we estimate 5\\%) escape detection, but we find a total gas content of the LA that is about 50\\% higher than previously assumed. These methods allow the deceleration of the LA clouds to be traced towards the Milky Way disk by drag forces. The derived velocity gradient strongly supports the assumption that the whole LA originates entirely in the Large Magellanic Cloud (LMC). LA4 is observed opposite to LA1, and we propose that both arms are related, spanning about 52\\,kpc in space. HT structures trace drag forces even at tens of kpc altitudes above the Milky Way disk.} ",
        "introduction": "Because of its proximity to the Milky Way, the MCS forms an ideal laboratory for studing the structure formation within the local universe in great detail. The more massive Milky Way galaxy accrete the MCS and tidal forces and perhaps drag separate gas from the stellar bodies. High--sensitivity optical studies disclose even the faintest stellar populations of evolved stars and allow to search for evidence of the stellar--gas--feedback processes triggered by the interaction with the Milky Way halo and/or its gravitational field \\citep{indu11}. Both stellar distributions are highly concentrated, so only a fraction also populate the Magellanic Bridge \\citep{putman03, irwin1990}. So far the prominent \\hi structures denoted as the Leading Arm \\citep[LA,][]{putman1998} and the Magellanic Stream have not detectable stellar counterpart \\citep[MS,][]{wannier1972, putman03, bruens05}.  Several investigations of the MS make use of specialized surveys \\citep{bruens05} or the Leiden-Argentine-Bonn all-sky survey \\citep{nidever08}. Even though a wealth of information has been compiled today \\citep[see][and references therein for a recent review]{nidever08}, it is still a matter of debate to what extent tidal or drag forces determine the structure formation of the MS and LA. Not detecting of stars in the MS argues for drag forces, decelerating only the viscous gaseous component, while the stellar component is unaffected and continues its orbital motion.  The existence of the LA, however still favors strongly a tidal origin \\citep{nidever08}.  In this paper we focus on the \\hi 21-cm emission distribution of the LA as observed with the Galactic All-Sky Survey \\citep[GASS,][]{naomi09,kalb10}. Assuming a linear distance of about 50\\,kpc from the Sun \\citep{cioni00}, single--dish radio telescopes like the Parkes 64-m dish offer a linear resolution of about 200\\,pc at \\hi 21-cm wavelength. The high spectral resolution of state-of-the-art radio spectrometers not only allows determination of the bulk motion of individual clouds but also examination of well resolved \\hi line profiles. GASS for the first time makes fully sampled sensitive wide field images of the \\hi distribution available.  For our study we extracted a huge portion of the southern sky from this data base to compile a comprehensive catalog of LA clouds and filaments (Sect.\\,\\ref{sec:data}). To achieve this aim we developed an approach to differentiating between Milky Way gas and superposed (in space and frequency) MCS \\hi emission. Based on a newly developed extension of a standard image recognition algorithm, we decompose the \\hi data without the necessity of Gaussian fits to the emission line profiles (Sect.\\,\\ref{sec:decompose}). This allows us to compile a comprehensive view of the LA down to a limit $T_{\\rm B} = 200$\\,mK. % Using this catalog we investigate the velocity structure of the LA and a newly identified leading arm four feature as a whole (Sect.\\,\\ref{sec:velostruc}). The catalog comprises the basic physical parameters of all identified LA objects; this gives us the opportunity to analyze the occurrence of head-tail structures and their orientation with respect to the MCS (Sect.\\,\\ref{sec:HT}). We end the paper with a brief summary (Sect.\\,\\ref{sec:summary}). ",
        "conclusions": "\\label{sec:summary} We used the GASS data \\citep{naomi09, kalb10} to investigate the LA structure. The analyzed area covers $70^\\circ \\times 85^\\circ$ degrees. Using a sophisticated image characterization tool, it was feasible to differentiate between the diffuse Milky Way emission and the LA structures. In total we identified 449 objects above the $4-\\sigma$ threshold of the original data ($T_{\\rm B} = 200\\,{\\rm mK}$). The applied method not only allows to determine the location of the objects, but also compiles derived information such as radial velocity, phase--space volume, column density, geometric center, and temperature weighted center. This allowed us to perform an automated statistical analysis of the LA as a whole and individual substructures down to the scales of individual clouds (HT structures).  Next to the well studied LA1, LA2, and LA3, we discovered LA4, which is localized at high Magellanic latitudes but at comparable longitudes to LA1. This symmetry but also the associated velocity gradient, suggests that this feature belongs to the LA complex. LA4 on average contains less extended clouds. In Sect. \\ref{sec:velostruc} we argued for a distance between 40 and 70 kpc, which implies that these clouds could also have a lower average mass.  Averaging across the whole LA1 to LA4 cataloged clouds, we derived a clear negative velocity gradient (deceleration) with Magellanic stream longitude. Assuming $D = 50$\\,kpc, the mass weighted linear correlation yields $v_{\\rm GSR} = 84.2\\,{\\rm km\\,s^{-1}}$ as axis intercept. This value is identical within its uncertainties with the bulk velocity of the LMC derived by \\citet{vandermarel02} from the carbon star ensemble. According to this we support the hypothesis of \\citet{nidever08} that the LA as a whole has its origin in the LMC.  The mass of the LA is about a factor of 50\\% higher than previously assumed (\\citet{bruens05} at a distance of 50\\,kpc, \\citet{cioni00}). The main reason for this increase is that our survey covered a significantly larger region than is available for previous, more specialized investigations. Some objects, blended with Milky Way emission features, may be missing in our analysis. The LA mass of $38 \\cdot 10^{6}\\,{\\rm M_\\odot}$ derived here is probably a lower limit only.  The velocity structure of individual LA filaments can be interpreted as prototypical for deceleration by drag forces. LA1.3 shows the strongest velocity gradient, consistent with the proposal of \\citet{naomi08}. All coherent structures of the LA but LA1.1 and LA1.2 show a systematic deceleration towards the Galactic plane. We find an obvious difference in the mass vs. velocity distribution for LA1 and LA4 in comparison to LA2 and LA3. The last two show that the higher the mass, the higher the $v_{\\rm GSR}$ velocity. In the framework of drag models, these high--mass compact objects keep their momentum, while the low--mass, more diffuse clouds are significantly decelerated.  The apparent symmetry of the LA structure is a surprising finding. The whole structure covers nearly $65^\\circ$ across the whole sky. At a distance of 50\\,kpc, this corresponds to a linear scale of 52.5\\,kpc. Combining the radial velocity information of the \\hi data with the proper motion \\citep{vandermarel02}, we estimate an inclination of only $13.6^\\circ$ against the plane of the sky. With the escape velocity for the LMC we can calculate that the LA structures are at least $6\\cdot 10^7$\\,a old. This nearly face-on geometry of the LA localized LA1.3 at a galactocentric distance of $R\\sim 19.3$\\,kpc and LA4 at $R\\sim 71.9$\\,kpc.  We find pronounced HT structures associated with LA4. Using a distance of 74\\,kpc, an estimate on the kinetic temperature by the \\hi line width, we estimate a pressure of $P = 60\\,{\\rm K\\,cm^{-3}}$. According to \\citet{kalb03} the ambient volume density with $n \\simeq 10^{-3}\\,{\\rm cm^{-3}}$ is high enough to cause drag.  The analysis of the ratio between $T_{\\rm B}$ and $T_{\\rm kin}$ implies that the gas of the objects is compressed towards the leading edge of the objects. Warm rims are at the head of the HT structures, suggesting gas compression and eventually an enhanced cooling. All massive LA objects show HT structures. The sky projected vector component between the location of their centers of mass relative to the maximum \\hi line intensities are all oriented parallel to the orbital motion of the MCS. In combination with the velocity segregation, the overall velocity gradient and the cool leading rims of the HT objects, we conclude that the whole LA is a coherent structure decelerated by drag. The newly identified LA\\,4 fits in the whole structure and physical interpretation perfectly.  The emerging picture of a four-arm structure, with two symmetric arms approaching the Milky Way disk (LA1 and LA4) and two arms that have already passed the disk (LA2 and LA3) provides new constraints for simulations of the orbital history of the Magellanic Clouds. \\citet{DiazBekki2011} present a new tidal model in which, for the first time, structures resemble a clear bifurcation of the LA. Concerning LA4, we would like to note that the existence of this feature was predicted by \\citet{Besla2010}, but their simulations fail to recover LA1. Hydrodynamic processes such as ram pressure are needed to shape the Magellanic Stream."
    },
    "1208/1208.4354_arXiv.txt": {
        "abstract": "We find that the viability of a cosmological model that incorporates 2 sterile neutrinos with masses around 1 eV each, as favored by global neutrino oscillation analyses including short baseline results, is significantly dependent on the choice of datasets included in the analysis and the ability to control the systematic uncertainties associated with these datasets. Our analysis includes a variety of cosmological probes including the cosmic microwave background (WMAP7+SPT), Hubble constant (HST), galaxy power spectrum (SDSS-DR7), and supernova distances (SDSS and Union2 compilations). In the joint observational analysis, our sterile neutrino model is equally favored as a $\\Lambda$CDM model when using the MLCS light curve fitter for the supernova measurements, and strongly disfavored by the data at $\\Delta\\chi^2_{\\rm eff} \\approx 18$ when using the SALT2 fitter. When excluding the supernova measurements, the sterile neutrino model is disfavored by the other datasets at $\\Delta\\chi^2_{\\rm eff} \\approx 12$, and at best becomes mildly disfavored at $\\Delta\\chi^2_{\\rm eff} \\approx 3$ when allowing for curvature, evolving dark energy, additional relativistic species, running of the spectral index, and freedom in the primordial helium abundance. No single additional parameter accounts for most of this effect. Therefore, if laboratory experiments continue to favor a scenario with roughly eV mass sterile neutrinos, and if this becomes decisively disfavored by cosmology, then a more exotic cosmological model than explored here may become necessary. ",
        "introduction": "\\label{introlabel2v}  The standard models of particle physics and cosmology do not yet fully describe the neutrino sector, with open questions related to the mass-generation mechanism of the neutrinos, any sterile neutrino partners of the active neutrinos, and their potential relation to the number of relativistic degrees of freedom inferred from cosmology. In recent years, there has been some experimental evidence pointing towards the existence of additional light (effectively massless) degrees of freedom. In particular, a combined analysis of cosmic microwave background (CMB) data from WMAP7, baryon acoustic oscillation (BAO) distances from SDSS+2dF, and Hubble constant from HST yields a weak preference for additional light degrees of freedom ($\\neff = 4.34 \\pm 0.87$)~\\cite{Komatsu:2010fb}. When moreover including small-scale CMB data from ACT or SPT, this preference mildly increases to the $2\\sigma$ level ($\\neff = 4.56 \\pm 0.75$ with addition of ACT~\\cite{dunkleyact} and $\\neff = 3.86 \\pm 0.42$ with addition of SPT~\\cite{Keisler:2011aw}). These constraints on $\\neff$ explicitly assume that the additional particles are massless, and have sparked further work~\\cite{Joudaki:2012fx,Hamann:2010bk,Archidiacono:2011gq,Riess:2011yx,Smith:2011es,Hou:2011ec, Hamann:2007pi, Hamann:2011hu, GonzalezMorales:2011ty, Calabrese:2011hg,Smith:2011ab, Hamann:2011ge,Giusarma:2011ex,Giusarma:2011zq,Fischler:2010xz,deHolanda:2010am,Nakayama:2010vs, Burenin:2012uy, Ciuffoli:2012yd}.  In light of new predictions for the anti-neutrino flux from nuclear reactors, global short-baseline neutrino oscillation data now favor the existence of two sterile neutrinos with best-fit masses of $m_4 = 0.68~{\\rm{eV}}$ and $m_5 = 0.94~{\\rm{eV}}$, assuming massless active neutrinos~\\cite{Kopp:2011qd} (also see~\\cite{Giunti:2011gz,Donini:2012tt,Abazajian:2012ys,Conrad:2012qt}). Instead of analyzing the data with the aim of estimating an upper bound to the mass of an additional thermalized neutrino species~\\cite{Hamann:2010bk,Giusarma:2011zq,Giusarma:2011ex}, we take the existence of two sterile neutrinos with $m_4$ and $m_5$ as a prior assumption consistent with the short-baseline data. It is our aim to determine how a model with these two additional neutrino species fares compared to the case without them, when including all available and relevant cosmological data.  \\begin{table}[t!] \\vspace{0.5em} \\begin{center} \\begin{tabular}{lc|c} \\hline\\hline Parameter & Symbol & Prior\\\\ \\hline Baryon density & $\\Omega_{b}h^2$ & $0.005 \\to 0.1$\\\\ Cold dark matter density & $\\Omega_{c}h^2$ & $0.01 \\to 0.99$\\\\ Angular size of sound horizon & $\\theta_s$ & $0.5 \\to 10$\\\\ Optical depth to reionization & $\\tau$ & $0.01 \\to 0.8$\\\\ Scalar spectral index & $n_{s}$ & $0.5 \\to 1.5$\\\\ Amplitude of scalar spectrum & $\\ln{(10^{10} A_{s})}$ & $2.7 \\to 4$\\\\ \\hline Effective number of neutrinos & \\neff &  $3.046 \\to 10$\\\\ {\\it~~~--~with sterile neutrinos} & \\neff &  $5.046 \\to 10$\\\\ Sum of neutrino masses & $\\sum{m_{\\nu}}~\\rm{[eV]}$ &  $0$\\\\ {\\it~~~--~with sterile neutrinos} & $\\sum{m_{\\nu}}~\\rm{[eV]}$ &  $1.62$\\\\ Constant dark energy EOS & $w$ & $-3 \\to 0$\\\\ Running of the spectral index & ${dn_s \\over d\\ln k}$ &  $-0.4 \\to 0.4$\\\\ Curvature of the universe & $\\Omega_{k}$ &  $-0.4 \\to 0.4$\\\\ Primordial helium abundance & $Y_p$ &  $0 \\to  1$\\\\ \\hline\\hline \\end{tabular} \\caption{We impose uniform priors on the above cosmological parameters. In addition, we always consider the Poisson point source power $D_{3000}^{\\rm{PS}}$, the clustered power $D_{3000}^{\\rm{CL}}$, and the SZ power $D_{3000}^{\\rm{SZ}}$ as nuisance parameters constrained by the CMB data~\\cite{Keisler:2011aw}. Moreover, we always derive $\\sigma_8$, the amplitude of linear matter fluctuations on scales of $8~{\\rm{Mpc}}/h$ at $z=0$. We only vary a redshift-independent dark energy equation of state (EOS). In this table, the first 6 parameters are defined as ``vanilla\" parameters. } \\vspace{-2.5em} \\label{table:priorsnu} \\end{center} \\end{table}  \\begin{figure*}[!t] \\begin{center} \\vspace{-0.4em} \\includegraphics[bb=4.540570 10.481554 672.656604 470.818040,clip,scale=0.35]{fig1.ps} \\includegraphics[bb=1.138570 3.064570 668.518081 476.020110,clip,scale=0.35]{fig2.ps} \\end{center} \\vspace{-1.9em} \\caption{ {\\it Left}: CMB temperature power spectrum measurements with WMAP7 (orange) and SPT (blue). The $\\Lambda$CDM model without sterile neutrinos is shown with the solid (black) line, and the $\\Lambda$CDM model with 2 sterile neutrinos is shown in dashed (red). {\\it Right}: Assuming the $\\Lambda$CDM model is centered on the DR7 data, with error bars given by the shaded band (in blue), we show the absolute difference with our sterile neutrino model in solid (red). } \\label{fig:cmbpknu} \\end{figure*}  We examine the impact of the two sterile neutrinos on other cosmological parameters in the vanilla $\\Lambda$CDM model, such as the matter density, amplitude of linear matter fluctuations on $8~{\\rm{Mpc}}/h$ scales, and spectral index. We also explore the impact of extensions of a cosmological model with sterile neutrinos, including nonzero curvature, evolving dark energy, running of the spectral index, and primordial helium abundance. Throughout this paper, we will assume that the two sterile neutrinos are thermally populated as seems reasonable given the large mixing angles of the sterile neutrinos~\\cite{DiBari:2001ua,Abazajian:2002bj}. If this is not the case, then the differences between a model with two sterile neutrinos and one without them will be smaller (cf.~Refs.~\\cite{Abazajian:2004aj,Hannestad:2012ky,Mirizzi:2012we}).  The cosmological influence of sterile neutrinos includes an increase in the effective number of neutrinos to $\\neff = 5.046$ and the sum of neutrino masses to $\\sumnu = 1.62~{\\rm{eV}}$ assuming full thermalization. As discussed in Ref.~\\cite{Joudaki:2012fx}, the effective number of neutrinos is mainly correlated with the matter density and spectral index in a vanilla $\\Lambda$CDM model. In extended cosmological models, correlations also exist with the helium abundance, dark energy equation of state, and running of the spectral index. Meanwhile, the sum of neutrino masses is mainly correlated with the matter density and Hubble constant in a vanilla $\\Lambda$CDM model, along with the dark energy equation of state and curvature density in extended parameter spaces~\\cite{Joudaki:2012fx}.  The radiation content of the universe can be constrained from big bang nucleosynthesis (BBN) through its effect on the expansion rate~\\cite{Kneller:2004jz,Simha:2008zj,Steigman:2007xt}. Given the standard BBN consistency relation between the set of parameters $\\{Y_p, \\neff, \\Omega_b h^2\\}$~\\cite{Simha:2008zj}, the inclusion of 2 additional neutrinos boosts the primordial helium abundance by $\\Delta Y_p = 0.024$ when the baryon density is kept fixed. Thus, $Y_p \\approx 0.27$ in standard cosmological analyses when enforcing this consistency relation. Primordial helium abundance estimations from observations of metal poor extragalactic H~II regions suffer from significant systematic uncertainties (e.g.~see~\\cite{Peimbert:2007vm,Izotov:2007ed,Izotov:2010ca,Aver:2010wq,Aver:2010wd,Aver:2011bw}). An extensive analysis that attempts to account for these systematic uncertainties gives $Y_p = 0.2534 \\pm 0.0083$~\\cite{Aver:2011bw}, which is consistent with the cosmological estimate at 95\\% CL (assuming 5 light neutrinos). This agreement could be tightened by lowering $Y_p$ from cosmology, achieved via mechanisms such as incomplete thermalization, presence of a non-zero chemical potential, or post-BBN production of the sterile neutrinos from the decay of a heavy particle species (e.g.~see~\\cite{Hamann:2011ge}).  We describe our analysis method in Section~2. In Section~3, we provide constraints on a $\\Lambda$CDM model with three massless active neutrinos and two massive sterile neutrinos, and determine how well this model fits cosmological data relative to a model without sterile neutrinos. We further explore to what extent the tension between the two models could be ameliorated by an extension of parameter space including evolving dark energy, universal curvature, running of the spectral index, additional relativistic species, and freedom in the primordial helium abundance (all parameters defined in Table~\\ref{table:priorsnu}). Section~4 concludes with a discussion of our findings. ",
        "conclusions": "Global short-baseline neutrino oscillation data seem to favor the existence of two sterile neutrinos with masses close to 1~eV each (assuming effectively massless active species). We have studied the extent to which these two neutrinos are allowed by a combination of probes including the cosmic microwave background, Hubble constant, galaxy power spectrum, and supernova distances. In the analysis of SN data, we considered the impact on our results of both the SALT2 and MLCS light curve fitters. In particular, we showed that the choice of the SN light curve fitting method has a major impact on the inferred cosmological model.  We find that the sterile neutrino model provides a good fit to each of the considered datasets, and no single probe manages to decisively disfavor the sterile neutrino model with respect to the null model. In the joint analysis, sterile neutrinos are allowed by the cosmological data ($\\Delta\\chi^2_{\\rm eff} \\approx 0$) when using the MLCS light curve fitter for the SNe in the SDSS compilation, and strongly disfavored by the data ($\\Delta\\chi^2_{\\rm eff} \\approx 18$) when using the SALT2 fitter for SNe in the Union2 compilation. When excluding the supernova measurements, the sterile neutrinos are disfavored by the other datasets at $\\Delta\\chi^2_{\\rm eff} \\approx 12$. For a 3+1 sterile neutrino model, it is conceivable that the tension is ameliorated, but this depends on the mass of the single sterile neutrino. As an illustrative comparison, a cosmological model (without sterile neutrinos) that has $w = -0.8$ is disfavored by WMAP+SPT+$P(k)$+HST (no SN data) at the $\\Delta\\chi^2_{\\rm eff} = 9.4$ level compared to the vanilla model with $w = -1$.  If the SALT2 fitter is indicative of the correct way to interpret SN light curve measurements, then reconciling two light ($\\sim \\rm eV$) sterile neutrinos (consistent with results from short-baseline neutrino oscillation data) with cosmology may require additional freedom in the cosmological model. However, no single parameter from among nonzero curvature, evolving dark energy, additional relativistic species, running of the spectral index, and primordial helium abundance was able to decrease $\\Delta\\chi^2_{\\rm eff}$ or $\\Delta {\\rm DIC}$ close to zero. In fact, even for an extended space with all of these additional parameters, the sterile neutrino model is mildly disfavored at $\\Delta\\chi^2_{\\rm eff} \\approx 3$ (when using the SALT2 fitter).  The important take-home message, however, is that large shifts in $\\Delta\\chi^2_{\\rm eff}$ ($\\sim 20$)  already occur from subtle changes to the way parts of the cosmological datasets are analyzed. If SN studies converge toward the MLCS fitter (as opposed to the SALT2 fitter), then two sterile neutrinos with masses close to the eV level are easily allowed by the data. Interestingly, even when assuming the existence of two massive sterile neutrinos, we continue to find about  $2\\sigma$ preference for an additional massless species. In addition, in this model with two sterile neutrinos, a much larger matter density would be required (by roughly 40\\%), which helps preserve the constraint on $\\sigma_8 (\\Omega_m/0.25)^{0.47}$ near the 0.8-mark, in agreement with galaxy cluster abundance measurements. The analysis presented in this paper shows that it is premature to either rule out the existence of two massive sterile neutrinos or claim this model is cosmologically preferred.   \\smallskip {\\it Acknowledgements:} We much appreciate useful discussions with John Beacom, Ryan Foley, Jan Hamann, Ryan Keisler, Jostein Kristiansen, Gregory Martinez, and Joseph Smidt. We acknowledge the use of CAMB and CosmoMC packages~\\cite{LCL,Lewis:2002ah}. KNA is partially supported by NSF CAREER Grant No.\\ 11-59224. MK and SJ were partly supported by NSF Grant No.\\ 0855462. This research was supported in part by the Perimeter Institute of Theoretical Physics during a visit by MK. Research at Perimeter Institute is supported by the Government of Canada through Industry Canada and by the Province of Ontario through the Ministry of Economic Development and Innovation."
    },
    "1208/1208.4024_arXiv.txt": {
        "abstract": "The scattering of $f-$modes by magnetic tubes is analyzed using three-dimensional numerical simulations. An $f-$mode wave packet is propagated through a solar atmosphere embedded with three different flux tube models which differ in radius and total magnetic flux. A quiet Sun simulation without a tube present is also performed as a reference. Waves are excited inside the flux tube and propagate along the field lines, and jacket modes are generated in the surroundings of the flux tube, carrying $40\\%$ as much energy as the tube modes. The resulting scattered wave is mainly an $f-$mode composed of a mixture of $m=0$ and $m=\\pm 1$ modes. The amplitude of the scattered wave approximately scales with the magnetic flux. A small amount of power is scattered into the $p_1-$mode. We have evaluated the absorption and phase shift from a Fourier-Hankel decomposition of the photospheric vertical velocities. They are compared with the results obtained from the emsemble average of 3400 small magnetic elements observed in high-resolution MDI Doppler datacubes. The comparison shows that the observed dependence of the phase shift with wavenumber can be matched reasonably well with the simulated flux tube model. The observed variation of the phase-shifts with the azimuthal order $m$ appears to depend on details of the ensemble averaging, including possible motions of the magnetic elements and asymmetrically shaped elements. ",
        "introduction": "The work by \\citet{Braun+etal1988} has shown that sunspots can absorb up to half the power of incident $p-$modes. At the same time, part of the wave flux is scattered. The $p-$mode absorption and scattering phase shifts depend on the frequency, degree, radial order, and azimuthal order of the incident mode \\citep{Bogdan+etal1993,Braun1995}, and also on the magnetic structure, making the study of the scattering a promising way to infer the subsurface structure of sunspots and other magnetic features. In these studies the authors used Hankel analysis, a method which decomposes the $p-$modes into inward and outward propagating waves in annuli surrounding the sunspot.  Several mechanisms have been proposed to explain the observed absorption, mode conversion \\citep{Cally+Bogdan1993} being the principal candidate. When an acoustic wave encounters a magnetic field concentration, it is split into fast and slow modes. This mode transformation occurs at the height where the Alfv\\'en and sound velocities are comparable, since at that layer the distinction between the modes is small. Below this height the sound speed is higher than the Alfv\\'en speed and the modes are effectively decoupled. The fast mode is an acoustic-like wave, while the slow mode is similar to an Alfv\\'en wave and propagates along field lines removing energy from the acoustic wave. The observed absorption of $f-$modes can be accounted for by a vertical magnetic field \\citep{Cally+etal1994}, but the $p-$mode absorption obtained from this model is insufficient. However, the presence of inclined magnetic field produces significant increases in absorption with a peak at around 30$^o$ \\citep{Crouch+Cally2003}, which are consistent with observed values \\citep{Cally+etal2003}. Other mechanisms may also play a role. One of them is resonant absorption \\citep{Hollweg1988, Rosenthal1992}. It may occur when the flux tube has a smooth variation of the magnetic field rather than discontinuous, and represents absorption of wave energy by the transition layer when the incident acoustic waves resonantly excite MHD waves in the magnetic structure. However, the amount of absorption achieved by this mechanism cannot explain the observed loss of $p-$mode power. \\citet{D'Silva1994} points out that apart from this absorption produced by the dissipation of the $p-$modes in resonant layers and mode conversion, mode mixing also takes places. In this process the power of an incident $p-$mode mode with a certain frequency and degree $l$ can be dispersed into an outgoing $p-$mode with the same frequency, but different degree.    Flux tubes are a key feature to understand solar magnetic activity. They are spread all over the solar surface and couple different layers of the solar atmosphere. It has been proposed that magnetic flux tubes can act as wave guides, being one of the possible sources which supply energy to the upper layers to account for the chromospheric and coronal heating \\citep{Jefferies+etal2006}. The interaction of $p-$modes with thin flux tubes excites tube waves, including sausage waves, which are axisymmetric, longitudinal waves driven by variations in the total pressure, and kink waves, whose restoring force is magnetic tension and buoyancy, producing transversal oscillations. These waves propagate upward or downward and extract energy from the $p-$modes of the acoustic cavity \\citep{Bogdan+etal1996, Hindman+Jain2008, Jain+etal2009}. The kink mode is driven by the distortion of the tube produced by the harmonic flow field of the $p-$modes, while the sausage wave is excited by the pressure pertubations associated to the acoustic waves \\citep{Bogdan+etal1996}. These mechanisms are different from mode conversion, discussed in the previous paragraph, where fast and slow magnetoacoustic waves exchange energy due to their strong coupling in the region where the sound and Alfv\\'en speeds are comparable. Recently, \\citet{Daiffallah+etal2011} used numerical simulations to study the scattering of an $f-$mode by vertical flux tubes of different sizes, finding that the scattering by tubes with small radius is dominated by the kink mode, while the sausage mode is dominant for large tubes. This result coincides with the earlier work by \\citet{Bogdan+etal1996}, who studied analytically the nature of wave interactions with thin flux tubes \\citep{Spruit1981} and found that the kink mode is the dominant tube wave. The thin flux tube approximation assumes that the diameter of the tube is smaller in comparison to the pressure scale height, and thus the horizontal variations inside it can be neglected. \\citet{Hanasoge+etal2008} evaluated the scattering matrix associated with an $f-$mode that interacts with a thin flux tube in a stratified atmosphere, focusing on the kink mode excited in the magnetic tube. They found that most of the scattered wave corresponds to an $f-$mode with amplitude of 1.17\\% and with a phase shift of abound 50$^o$ relative to the incident wave, overstimating the observed value by a factor of 8.8 \\citep{Duvall+etal2006}. On the other hand, the recent work by \\citet{Hindman+Jain2011} analyzed the axisymmetric scattering of $p-$modes, mediated through the excitation of sausage waves on the flux tube, instead of the kink mode. They obtained a small absorption due to the poor coupling between the $f-$mode and the sausage mode for thin flux tubes, as pointed out by some of the works previously described in this paragraph.  Although these theoretical works have provided the first predictions about the modification of the solar wave field produced by flux tubes, as far as our concern no attempt has been made to observationally measure the detailed properties of the scattering produced by these small magnetic elements, with the exception of the estimates of amplitude and phase of monopole and dipole scattering by \\citet{Duvall+etal2006}. One of the objectives of this work is to present the measurement of the phase shift and its variations with the azimuthal orders $m$ and degree $L$. This data is a fundamental input to confront with the theory. On the other hand, most theoretical studies of this topic have been based on an analytical development. All of them have been restricted by some limitations, including the use of the thin flux tube approximation, the lack of the gravitational stratification, the analysis of a polytrope instead of a realistic solar atmosphere, or some constrains in the process that mediates the scattering. Numerical simulations are a more versatile approach and allow us to study more general situations.   In this work we study the scattering of an $f-$mode by flux tubes of different radius and magnetic flux using numerical simulations. As discussed previously, it is well known that thin flux tubes support sausage and kink modes \\citep{Roberts+Webb1978, Spruit1981}. On the other hand, in unstratified atmospheres permeated by homogeneous magnetic fields one would expect the propagation of pure fast and slow magnetoacoustic waves and Alfv\\'en mode. In atmosphere stratified by gravity (for example) the fast, slow, and Alfv\\'en waves are coupled in general and this distinction between modes no longer applies, but even in these cases it is useful to refer to this simple picture to discuss the properties of fast and slow magnetoacoustic-like waves (in regions where the sound and Alfv\\'en speed differ greatly). In the case of thick flux tube models, representative of a sunspot, for example, acoustic waves can be converted into these modes by means of mode conversion. The flux tube models presented in this paper correspond to an intermediate case between these two extremes. We expect a smooth transition (with increasing radius) from mostly excitation of the kink and sausage modes at small radius, to excitation of waves that look like the fast and slow magnetoacoustic waves of a thick flux tube (at large radius). However, in this study we made no attempt to distinguish between neither the different wave modes which are present nor the mechanisms that generates them. Instead, we will refer to ``tube modes'' or ``tube waves excitation'' indistinctly. We aim to carry out a direct comparison between the numerical and observational results by performing a Hankel analysis of the data obtained from the interaction of an $f-$mode with flux tubes. The organization of the paper is as follows. In Section \\ref{sect:observations} we introduce the observations used in this study. Section \\ref{sect:procedures} briefly describes the numerical code and the set up of the simulations. In Sections \\ref{sect:scattering} and \\ref{sect:jacket}, we present the tube mode excitation, the scattering, and the jacket modes produced in these simulations. The results of the Hankel analysis are shown in Section \\ref{sect:hankel}, including a comparison with observations, and finally we conclude with a summary of our calculations and a discussion of their applicability to understanding observations. ",
        "conclusions": "\\label{sect:conclusions}  We have presented the analysis of the scattering produced by magnetic flux tube models using 3D numerical simulations. Previous attemps to model this phenomenon \\citep[\\eg,][]{Gordovskyy+Jain2007, Jain+etal2009,Hanasoge+etal2008, Hindman+Jain2011} have faced the problem by means of analytical treatments. These types of studies are the first steps toward the comprenhension of this issue and they provide a valuable heritage to understand the wave interaction with magnetic media and confront it with the forthcoming observations and modeling. However, the simplifications needed to carry out their development restrict their results to some idealized cases. From this scope, the use of numerical simulations emerges naturally as the next step to address these questions in more general situations.  In these simulations we have propagated an $f-$mode through a model S atmosphere \\citep{Christensen-Dalsgaard+etal1996} stabilized against convective instabilities embedded with a flux tube model. In order to compare how some properties of the tubes affect the scattering three realizations were performed, using different flux tube models. All tubes have the same peak magnetic field strength, but they differ in the radius and, thus, in the magnetic flux.  Our simulations show that the interaction of an $f-$mode with a flux tube excites tube waves. These waves propagate along the magnetic field lines and produce a real absorption of the incident energy, since it is extracted from the acoustic cavity. The oscillations of the tube produced by these waves generate a scattered wave. It is composed of a mixture of axisymmetric ($m=0$) and dipolar ($m=\\pm1$) modes, whose distribution in frequency and azimuthal order depends on the radius of the flux tube. For thin flux tubes, the $m=\\pm1$ dipolar oscillation dominates the tube wave, while axisymmetric oscillations ($m=0$) become important for larger tubes. This result agrees with those previously obtained by \\citet{Daiffallah+etal2011}.  We have quantified the absorption coefficient and phase shift produced by the three magnetic flux tube models. Based on the results discussed in the previous section, we draw the following conclusions. Firstly, the absorption increases with wavenumber (frequency) for all azimuthal orders and tube models. Secondly, the amount of absorption in general increases with the magnetic flux of the tube, although this increase depends significantly on the wavenumber and azimuthal order. Thirdly, the distribution of the absorption in azimuthal order depends on the radius of the tube. In all models the peak absorption is obtained for $m=1$. However, in the tube with 170 km radius it is followed by $m=0$, with a weak absorption in $m=2$, while in the 370 km radius tube the absorption in $m=2$ is stronger than the corresponding to $m=0$, and in the case of the 560 km radius tube the absorption in $m=0$ is especially low. The absorption of the axisymmetric $m=0$ order decreases with the radius.  It is noticeable that the different behavior that the phase shift shows regarding the second and third points of the previous paragraph. From the simulations with the tubes we find a similar distribution in the phase shift produced in different azimuthal orders, although it seems to approximately scale with the magnetic flux, with some dependence on the $L$ value and the azimuthal order. In this way, $m=0$ and $|m|=1$ show a very similar phase shift, the later slightly higher in the two smaller tubes and the opposite in the larger tube, while the phase shift produced in $m=2$ is very small, except for the 560 km radius tube.  In this work we are not only interested in modeling the scattering process, but also in applying this knowledge to interpret observations. A deeper understanding of the wave interaction with small magnetic scatterers can yield a basis to infer the properties of the scattering elements, even at scales smaller than the observational resolution. We have compared the numerical results for the phase shift with observations of an ensemble averaging of thousands of small magnetic elements. In order to perform an equivalent comparison, the phase shift obtained from 3400 realizations of the Hankel analysis of the simulations with the 560 km radius tube with a small shift in the position of the annulus was also averaged.    The phase shift produced by our larger tube model after averaging shows a good qualitative agreement with the observed phase shift. Since the phase shift scales with the magnetic flux of the scattering element, we may consider that the phase shift of the observed elements could be produced by fluxtubes with magnetic flux around the corresponding to tubes with 560 km radius and 1600 G. The current work suggests one possible solution for the properties of the tube model, although other combinations of radius and magnetic field strength might also work. This kind of measurement seems to be a promising method to infer the characteristics of small magnetic networks elements. However, some caution must be considered in their interpretation. Some of the observed magnetic elements used in this study show strong asymmetries (see Figure \\ref{fig:magnetogram}). When the scattering element is not axisymmetric the scattering is not restricted to occur from an azimuthal order $m$ to the same order $m$, but the scatterred wave can correspond to a different order $m'$. These nonsymmetric elements could contribute to the broadening of the distribution of phase shift with $m$. On the other hand, in the observed magnetograms it is hard to find enough completely isolated magnetic elements. From the 3400 elements used in the analysis, many of them present other small magnetic features inside the 16.5 Mm annulus in which the Hankel decomposition was performed. The observational analysis might be contaminated by the scattering produced by these other elements.  In the comparison between the observations and the simulations, we have assumed that the observational shift in the center of the Hankel analysis is restricted to a Gaussian with the FWHM of the PSF from high-resolution MDI. If the proper motions of the flux tube have larger extension, the results could be affected.  However, the tubes are presumably moving at the time scale of the granulation, which is not so different from the time scale of the wave period. This could mean that the approximation of a stationary tube is not so good, and limits the capacity of this work to address this issue. It would be interesting to extend the analysis to moving flux tubes in the future.  These results provide a warning to be cautious when interpreting ensemble averages of observational data. In observations like those presented in this paper, where the contribution of an individual magnetic feature is too low to get a reliable measurement, the average of several cases is a compulsory procedure to obtain a strong enough signal. However, the individual and unique characteristics of each element, together with the limitations to perform the analysis using exactly the same configuration, can lead to a result which may point to conclusions that do not reflect the real observed structure. In the particular case studied in this work, from the observational broad distribution of the phase shift with $m$, one could assume that a big magnetic feature is necessary to produce that dependence with the azimuthal order. The analysis of the averaged simulated flux tube reveals that a similar measure of the phase shift dependence can be retrieved from a very different magnetic element, making difficult infering an irrefutable conclusion about the nature of the observed elements."
    },
    "1208/1208.4212_arXiv.txt": {
        "abstract": "We investigate the possibility of Photospheric Radius Expansion (PRE) during magnetar bursts. Identification of PRE would enable a determination of the magnetic Eddington limit (which depends on field strength and neutron star mass and radius), and shed light on the burst mechanism.  To do this we model hydrostatic atmospheres in a strong radial magnetic field, determining both their maximum extent and photospheric temperatures.  We find that spatially-extended atmospheres cannot exist in such a field configuration:  typical maximum extent for magnetar-strength fields is $\\sim$ 10m (as compared to 200 km in the non-magnetic case).  Achieving balance of gravitational and radiative forces over a large range of radii, which is critical to the existence of extended atmospheres, is rendered impossible in strong fields due to the dependence of opacities on temperature and field strength.   We conclude that high luminosity bursts in magnetars do not lead to expansion and cooling of the photosphere, as in the non-magnetic case.  We also find the maximum luminosity that can propagate through a hydrostatic magnetar atmosphere to be lower than previous estimates.   The proximity and small extent of the photospheres associated with the two different polarization modes also calls into question the interpretation of two blackbody fits to magnetar burst spectra as being due to extended photospheres. ",
        "introduction": "Photospheric Radius Expansion (PRE) events can occur during bursts on neutron stars when the luminosity of the object reaches the Eddington Luminosity, i.e. where the radiation force balances the gravitational one: \\begin{equation} \\label{eq:Ledd} L_{\\rm Edd} \\equiv \\frac{4\\pi GM_*c}{\\kappa_\\rmn{Th}}, \\end{equation} (where $M_*$ is the stellar mass and $\\kappa_\\rmn{Th}$ the Thomson scattering opacity) and the large radiation pressure forces the atmosphere to expand outwards. For the hydrogen atmosphere of a $1.4 M_\\odot$ neutron star, $L_{\\rm Edd} = 1.8\\times 10^{38} \\rm{erg}\\rm{s}^{-1}$, while for a helium atmosphere it is twice that value. As a result, the photosphere moves to a much larger radius, corresponding to a drop in temperature $T$. For a neutron star with a modest magnetic field (up to $\\sim10^{12}~\\rm{G}$), Compton scattering dominates the opacity in the atmosphere and various relativistic effects allow the atmosphere to expand up to hundred kilometres, so that the temperature of the expanded photosphere drops out of the X-ray range altogether \\citep{Hoffman1978, Paczynski1986}. The hallmark of PRE in neutron stars is thus a `double-peaked' structure in the X-ray light curve of a burst, in which the flux increases to a maximum and then drops sharply (indicating the black-body temperature has decreased), before rising again steeply to a slightly larger maximum as the bolometric luminosity drops again and the photosphere contracts \\citep{Paczynski1983}.  PRE is characteristically seen in Type I X-ray bursts from accreting neutron stars, in which the build-up of accreted material leads to a thermonuclear explosion on the surface of the star, causing a huge increase in luminosity. PRE bursts have typically been used to constrain the mass and radius of the neutron star, thus potentially constraining the equation of state of dense matter \\citep[e.g.][]{Damen1990, Galloway2003a, Ozel2009a, Ozel2010a, Steiner2010a, Suleimanov2011a}. However, PRE is generically driven by high luminosities, irrespective of the underlying energy source. This has led to the recent suggestion that it might also happen in bright bursts from magnetars \\citep{Watts2010} -- isolated neutron stars with dipole magnetic fields above $\\sim10^{13}~\\rm{G}$ -- whose bursts (which occur over a wide range of luminosities) are thought to be powered by large-scale reconfiguration of the decaying magnetic field \\citep{Thompson1995}. \\cite{Watts2010} argued that observing PRE in magnetar bursts could put interesting constraints on the emission mechanism, magnetic field strength and mass-radius relationship for magnetars.  The suggestion that PRE might happen in magnetar bursts was motivated by the August 2008 observation of a large ($L_X \\sim 7\\times 10^{39}$ erg s$^{-1}$) burst from SGR 0501+4516, which showed a double-peaked light curve similar to those seen in Type I X-ray bursts. In their paper on this burst, \\cite{Watts2010} laid out several criteria required for PRE to occur, and argued that they were in general met for magnetar bursts. For PRE to occur in a neutron star, the flux must be emitted from an optically thick region, the radiation pressure must be sufficient to overcome gravity and other confining forces, the emitting region must remain optically thick during the expansion (so that the emission remains close to blackbody and effective temperature decreases with increasing photosphere radius), and the opacity must increase with distance from the star. The last point is slightly subtle: in order for the photosphere to expand, the luminosity $L$ must remain close to the critical luminosity ($L_{\\rm cr}$)  needed to balance radiation pressure with the confining forces\\footnote{Here and throughout the paper, we follow \\cite{Paczynski1986} in defining $L_{\\rm cr}$ as the actual maximum luminosity as a function of radius, modified by the changing opacity and gravitational redshift, whereas $L_{\\rm Edd}$ is strictly given by eq \\ref{eq:Ledd}}. However, both these quantities are modified differently by the strong gravitational field, so that $L/L_{\\rm cr} \\propto 1+z$, where $z$ is the gravitational redshift. To ensure that this quantity does not decrease with radius, which would make expansion impossible \\citep{Paczynski1986}, the opacity (which determines $L_{\\rm cr}$) must therefore increase with radius. In Type I X-ray bursts, this is effected by the Klein-Nishina corrections which reduce the Thomson cross-section at high temperatures close to the stellar surface.  The presence of a magnetar-strength magnetic field complicates the hydrostatic expansion of the atmosphere in three significant ways. In the  `trapped fireball' picture of magnetar bursts \\citep{Thompson1995}, the huge release of magnetic energy leads to the creation of a pair-plasma, so that the atmosphere is dominated by this dense pair-gas, rather than baryonic matter ablated from the star's surface. Further, closed field lines provide a strong confining force for both baryonic and leptonic matter, since plasma cannot easily move perpendicular to the field. A straightforward calculation (see section \\ref{sec:Lcr}) demonstrates that a very strong field can easily confine even the largest giant flares with $L\\sim10^{44} \\rm{erg}; \\rm{s}^{-1}$ \\citep{Lamb1982}, so that PRE in a magnetar will likely only occur in open field line regions.  The final effect of the magnetic field, and the most important one in the present work,  is to modify the electron scattering cross-sections by several orders of magnitude, depending on the polarization state of the scattering photon. The strong magnetic field suppresses electron motion perpendicular to the field, so that photons that try to excite this motion (i.e. that are polarized perpendicular to B; the `Extraordinary' or E-mode) have a greatly reduced scattering cross-section compared to the Thomson scattering cross section $\\sigma_\\rmn{Th}$, while photons polarized parallel to B (the `Ordinary', O-mode) are largely unaffected. The modified polarization-dependent cross-sections will increase the critical luminosity (sometimes called the `magnetic Eddington limit') by several orders of magnitude \\citep{Thompson1995, Miller1995}. Additionally, since the E-mode cross-section scales roughly as  $(T/B)^{2}$ (see Equation \\ref{eq:crosssec_red}), the opacity increases steeply with distance from the star from the decrease in field strength, and becomes strongly temperature dependent. $L_{\\rm cr}$ is therefore a strong function of radius. As we will demonstrate, the strong temperature and field dependence of the opacity has a profound effect on the structure of the magnetar atmosphere in comparison to the non-magnetic case.  The object of this paper is to explore the structure of a magnetar atmosphere at very high luminosities. We follow the approach of \\cite{Paczynski1986} and calculate the structure of a series of hydrostatic atmospheres with different masses, base temperatures, and magnetic field strengths, solving the equations for stellar structure. The main difference from \\cite{Paczynski1986} is in our consideration of the opacity, which is dominated by electron scattering in both cases. They consider a non-magnetic star, for which the electron scattering cross-section is the Thomson one (modified at high temperatures by Klein-Nishina corrections). We instead use the cross-sections modified by the strong magnetic field, so that the radiation is split into two polarization modes, and only E-mode photons diffuse through the atmosphere.  The paper runs as follows. In Sections \\ref{sec:Lcr}-\\ref{sec:structure} we discuss the concept of a `critical luminosity' in more depth, and present the equations we use to calculate the structure of the atmosphere. Section \\ref{sec:opacity} focuses on the electron scattering cross-sections in a super-strong magnetic field. The field introduces several complications (such as the dependence on photon energy of the scattering cross-section and the presence of the cyclotron resonance), and we explain how we calculate the effective opacity for a thermal photon distribution. Section \\ref{sec:results} presents the main results of our calculations, and Section \\ref{sec:comparison} compares our results to previous calculations of the critical luminosity. Finally, in Section \\ref{sec:extensions}  we examine additional physical processes (most significantly, a pair plasma gas and magnetic confinement from closed field lines), and argue that neither of these are likely to affect the qualitative conclusions of the paper, by which we mean the non-existence of hydrostatic extended magnetar atmospheres. ",
        "conclusions": "\\label{sec:conclusions}  Hydrostatic extended magnetar atmospheres do not exist, due to the fact that the strong dependence of opacity on radius makes the required near equality of the luminosity and the critical luminosity throughout the atmosphere impossible. A hypothetical extended magnetar atmosphere will thus either have large regions where luminosity is greater than the critical luminosity, or it will have large regions where the luminosity is much lower than the critical luminosity, so that it is not supported against gravity. Therefore, such an atmosphere cannot exist. This is unlike the nonmagnetic case, where a precise balance between luminosity and critical luminosity means the photospheric radius can extend up to 200 km. The fact that magnetars do not have extended hydrostatic atmospheres means that PRE, as envisaged previously, cannot work, due to the fact that stable expansion of the atmosphere is not possible.  Additionally, we find that the maximum luminosity that can be propagated through a hydrostatic magnetar atmosphere may be lower than the critical luminosity given by \\citet{Miller1995}, depending on the structure of the atmosphere. This is due to the fact that \\citet{Miller1995} assumes the maximum luminosity that can propagate will be set by the scattering of O-mode photons near the O-mode photosphere. However, in our models the maximum luminosity is set by E-mode photons scattering in the highest temperature region near the surface of the star, where the scattering cross section is relatively large due to the high frequency of the thermalized photons. This means that depending on the atmospheric structure, more observed magnetar bursts might have reached their critical luminosity than previously assumed.  Our results have implications for interpretation of spectral fits to magnetar burst data. Magnetar bursts are typically fit with several different spectral models, the main ones being two power laws, a power law plus a black body, two black bodies, and more recently a power law with exponential cutoff and optically thin thermal bremsstrahlung. The two black body model is typically one of the best fitting of these models \\citep{Olive2004a, Feroci2004a}, with typical fitting parameters giving two temperatures around 3 and 11 keV, and typical emission region radius of roughly the radius of a neutron star for the colder black body, and an order of magnitude smaller than that for the hotter component, but with a large scatter in the sizes.  While this model has mostly been presented as purely phenomenological, it has occasionally been suggested that this could be interpreted as representing the distinct signatures of the E- and O-mode photospheres \\citep[][]{Israel2008a, Kumar2010a}. This interpretation was attractive in the sense that it opened up the possibility of another way of measuring the critical luminosity. Our results show that in a hydrostatic atmosphere the temperature difference between these two photospheres is never more than one or two keV, and that they are very close to each other spatially. This makes the interpretation of the two blackbody model as representing the E- and O-mode photospheres highly unlikely.  Our results pose the question of what does happen when a magnetar burst exceeds the critical luminosity. As we have shown, hydrostatic extended atmospheres are impossible. Thus, solving this problem will require dynamical, time dependent models. These models would very likely result in outflows, which could give rise to several forms of observable emission, such as the radio emission seen from the outflow from the SGR 1806-20 Giant Flare \\citep{Cameron2005a, Gaensler2005, Granot2006}, or the `Magnetar Wind Nebula' recently detected around Swift J1834.9-0846 \\citep{Younes2012a}. Any outflows would likely be very sudden, due to the short sound crossing time and burst time scale involved. The observation that prompted this research, that of SGR J0501+4516,  likely reached or exceeded its critical luminosity \\citep{Watts2010}, and would thus be a prime candidate for testing any future dynamical models of magnetar bursts."
    },
    "1208/1208.0211_arXiv.txt": {
        "abstract": "{We present 3-dimensional kinematical observations of the massive OB association Cygnus~OB2 to identify the mechanisms responsible for disrupting young star clusters. The picture revealed by these observations is of a highly-substructured, dynamically unmixed OB association that does not exhibit the position--velocity correlations predicted by the theories of infant mortality or tidal stripping. These observations would appear to support a picture of hierarchical star formation.} ",
        "introduction": "It has been known for many years that there is a lack of old clusters compared to an extrapolation of the young cluster population. This was first recognized by Oort [17] and exacerbated by the discovery of large numbers of embedded clusters in the near-IR [13]. Based on this it became clear that the vast majority (at least 90\\%) of clusters disperse within 10~Myr [14].  This is usually explained by the process of {\\it infant mortality}, whereby residual gas left over from star formation is forced out of the cluster via feedback from massive stars, leaving the stellar part of the system in a super-virial state and prone to dissolution. This long-established theoretical framework assumes that the cluster was in virial equilibrium prior to gas expulsion and identifies the star formation efficiency (the fraction of gas turned into stars, $\\sim$5-10\\%) as the dominant factor in determining cluster stability [10,1]. This has recently been called into question: numerous theoretical works have suggested other parameters of equal importance such as the spatial distribution of stars at birth [15], the rate of residual gas expulsion [3], and the initial virial state of the cluster [16]. Overall, these studies suggest that the influence of gas expulsion may have been over-estimated and that many clusters may be stable from a young age.  An alternative theory, first put forward by Spitzer [18] to explain the lack of Gyr-old clusters, is that clusters are tidally heated by passing interstellar clouds. This has recently been used to explain the disruption of very young clusters by their parental giant molecular cloud (GMC) [7,12]. It has been argued to be an effective disruption mechanism for young clusters since the average cloud density in GMCs exceeds the tidal density from the galactic potential by an order of magnitude or more [7].  A third and important explanation for the lack of mature, gravitationally-bound clusters is that the majority of young clusters are not gravitationally bound, one of the fundamental assumptions of the theory of infant mortality. The discovery of vast numbers of near-IR embedded clusters led many authors to conclude that all stars formed in bound clusters [13,6,14]. However many of these clusters may just be stellar overdensities and may not be gravitationally bound. In their study of the spatial distribution of young stars, Bressert et al. [5] could find no evidence for a preferred scale for clustering that would be apparent if stars preferentially form in clusters (see also [8]). These findings support the view that star formation is hierarchical with no preference for scale and therefore most stars may not form in bound groups [2]. If the majority of stars do not form in bound clusters, then an efficient disruption mechanism may not be required.  To answer the question of how star clusters are disrupted, and thereby also address the question of whether all stars form in clusters, we must study a cluster in the process of being disrupted. This is difficult because the majority of clusters are either still embedded in their GMC (e.g. the Orion Nebular Cluster) or if they have already removed their residual gas they are often found to be gravitationally bound. This intrinsic bias is because clusters that display a clear overdensity and a spherical shape but which have also emerged from their parental GMC {\\it must} be bound if they have retained their clustered morphology. To study a cluster in the act of dispersal and therefore probe the mechanisms responsible we should not study star clusters but instead study OB associations, less dense groups of stars that have been suggested to be the result of expanded clusters [13]. ",
        "conclusions": "We present 3-dimensional dynamical observations of Cyg~OB2 to elucidate the physical mechanism that led to the association being gravitationally unbound. The overall picture of these observations is a highly substructured, dynamically unmixed association that does not exhibit the position--velocity correlations expected for the theories of infant mortality or tidal stripping. These observations appear to support a picture of hierarchical star formation, in which stars are not born in dense clusters, but in looser associations that retain substructure as they dynamically evolve. Further observations in Cyg~OB2 and other regions are necessary to confirm this."
    },
    "1208/1208.2813_arXiv.txt": {
        "abstract": "The usefulness of angle-action variables in galaxy dynamics is well known, but their use is limited due to the difficulty of their calculation in realistic galaxy potentials. Here we present a method for estimating angle-action variables in a realistic Milky Way axisymmetric potential by locally fitting a St\\\"ackel potential over the region an orbit probes. The quality of the method is assessed by comparison with other known methods for estimating angle-action variables of a range of disc and halo-type orbits. We conclude by projecting the Geneva-Copenhagen survey into angle-action space. ",
        "introduction": "In the study of dynamical systems it is becoming increasingly important to be able to process and understand large multi-dimensional data sets efficiently. The stars in our own galaxy, the Milky Way, are being increasingly observed in full six-dimensional phase-space through the combination of astrometry and radial velocity measurements. Full 6D phase-space information is currently available for stars in the solar neighbourhood from the Geneva-Copenhagen and RAVE surveys \\citep{Nordstrom2004,Zwitter2008} and this is to be greatly expanded on by the future space mission {\\it Gaia} \\citep{Perryman2001}. Beyond our Galaxy, the advent of integral-field spectroscopy has led to projects such as SAURON \\citep[and subsequent papers]{Bacon2001}, which mapped the kinematics of a representative sample of 72 nearby elliptical and spiral galaxies, and subsequently ATLAS3D \\citep[and subsequent papers]{Cappellari2011}, which combined SAURON observations with CO and HI observations to study the kinematics of a complete volume-limited sample of 260 local early-type galaxies. Observational data is often understood by performing large N-body simulations. Whilst such models are straightforward to produce, the configurations of the models are difficult to control and characterise. Schwarzschild modelling offers an improvement on this by describing the configuration of a model by a weighted set of orbits. However, this approach is not the most natural as each orbit is characterised by its initial phase-space coordinates. It is necessary that techniques are developed which can simplify both observational and simulation data without losing the richness of the phase-space information.   Angle-action variables are a set of canonical coordinates which can be used to express the equations of motion in a trivial form: the actions are integrals of the motion whilst the angles increase linearly with time. Such a formulation instantly reduces the complexity of any dynamical data set by reducing the six phase-space dimensions to three angle coordinates. Angle-action variables can be defined for any quasi-periodic orbit. Initially introduced to study celestial mechanics, angle-action variables now have great potential for galaxy dynamics due to their attractive properties. For instance, the Jeans theorem states that the arguments for the distribution function of a steady-state galaxy must be integrals of the motion, and it is particularly convenient to use as integrals the actions as \\begin{inparaenum}[(i)] \\item they are adiabatic invariants, \\item the zero-point of an action is well defined and \\item the range of values an action may take is independent of the other actions. \\end{inparaenum} The angle-action variables also provide a basis for the development of a perturbative solution to the equations of motion (see \\citealt{BT08} for a much fuller discussion of the merits of angle-action variables).  The increasing evidence of substructure within the stellar halo of the Milky Way \\citep[e.g.][]{Belokurov2006} has led many authors to consider the use of angle-action variables when hunting for and understanding the formation of structure within phase-mixed data sets. For example, \\citet{McMillanBinney2008} studied a simulation of a self-gravitating satellite in a realistic Galaxy potential in angle-action space. Though the stars became well phase-mixed, the action space still showed considerable structure and through the use of angle-variable diagnostics the Galaxy potential and history of the satellite could be reconstructed. Similarly, \\citet{Sellwood2010} and \\citet{McMillan2011} have used the Geneva-Copenhagen survey to analyse stars in the solar neighbourhood in angle space and showed that the Hyades moving group may be due to a recent inner or outer Lindblad resonance. The study of tidal streams has a natural expression in angle-action variables. Under certain conditions the dimensionality of the stream may be reduced to one as the stream stretches out in a single angle coordinate \\citep{Tremaine1999}. \\citet{EyreBinney2011} showed that the path of a stream can be reconstructed far more reliably in angle-action space than by incorrectly assuming that streams delineate the orbits of their progenitors.    Despite the aforementioned advantages, angle-action variables remain awkward to work with in practical applications due to the difficulty of their calculation in a general potential. They are easily calculated when the potential is spherical and with more work can be analytically calculated when the potential is of St\\\"ackel form, but neither of these approaches is satisfactory when working with realistic galaxy potentials, because such potentials do not satisfy these conditions. The development of methods to estimate angle-action variables in a general potential is crucial if we are to benefit from the advantages of angle-action variables and the wealth of techniques which utilise them.  In this paper we present a method for estimating angle-action variables in a general axisymmetric potential. The method proceeds by fitting a St\\\"ackel potential locally to the region of the potential a given orbit explores, thus enabling us to calculate analytically the actions and angles in this fitted St\\\"ackel potential. In Section~\\ref{AAinStackelPot} we give a brief overview of the determination of angle-action variables in an axisymmetric St\\\"ackel potential and then in Section~\\ref{StackFit} present the method for locally fitting such a potential to any axisymmetric potential. The results of the method are examined by analysing artificial data in Section~\\ref{Application} and then these results are compared to other methods in Section~\\ref{OtherMethodComparison}. Finally we demonstrate the practical application of the method by inspecting the Geneva-Copenhagen Survey in angle-action space in Section~\\ref{GCS}. ",
        "conclusions": "We have detailed a method for estimating the angle-action variables in a general axisymmetric potential given a 6D phase-space point. The method is based on locally fitting a St\\\"ackel potential to the region of the potential the orbit probes and then taking advantage of the ease with which we may calculate the actions and angles in a St\\\"ackel potential. We have investigated the systematic errors by producing phase-space points of known actions and angles using the torus machine and then assessing how well the method can reproduce these variables. For a single torus the errors in the angles are largest for phase-space points near apsis and the errors in the actions are constant (of order a few percent). For a collection of tori, chosen to be representative of both disc and halo-type tori, the absolute error in the actions is found to scale with the sum of the vertical and radial actions. The errors in the angles scale with the relative error in their corresponding action. We compared the method to other methods for estimating the actions in an axisymmetric potential. The method gives results approximately two orders of magnitude more accurate than assuming the potential is spherical and performs approximately a single order of magnitude better than the adiabatic approximation.  We have demonstrated that the procedure is suitable for most disc and halo-type orbits. The procedure will not work for resonant orbits or chaotic orbits. However, the occurrence of these orbits in realistic galaxy axisymmetric potentials is rare and the great majority of stars are on quasi-periodic non-resonant orbits \\citep{Ollongren1962,MartinetMayer1975}.  We demonstrated the use of the method by application to the Geneva-Copenhagen Survey (GCS). As this survey is only local, the angle-action space does not reveal much more information than velocity space. However, we present angle-action coordinates for the peaks of the clumps and streams present in the survey and use them to study the relative impact on estimated angles and actions of observational errors and the known systematic errors of the method. We show that the observational errors are dominant.  It is hoped that this method will lead to more widespread use of angle-action variables when analysing data, and we intend to release the source code soon\\footnote{The source code will be made available at http://galaxies-code.physics.ox.ac.uk.}. Whilst the GCS can be easily analysed in velocity space, angle-action variables should enable us to reveal structures, which are more dispersed in phase-space, in larger surveys.  We have limited the discussion in this paper to axisymmetric potentials. Whilst for our own spiral galaxy the axisymmetric approach may suffice, for analysing elliptical galaxies a triaxial approach must be developed. There are also triaxial St\\\"ackel potentials \\citep[see][]{deZeeuw1985a} so it should be possible to expand the approach outlined in this paper to triaxial potentials. The extension of the angle-action estimation is simple, but the fitting procedure is more complex when the potential is triaxial. \\citet{deZeeuw1985b} discuss how a general triaxial potential may be fitted both locally and globally by a St\\\"ackel potential. The method for global fitting is the three-dimensional generalisation of the method used in this paper so involves multiple multi-dimensional integrals. Also the best choice of coordinate system involves minimising the least-square difference with respect to two coordinate parameters so a more computationally expensive procedure than the simple method used in this paper may be required for finding the best coordinate system."
    },
    "1208/1208.5869_arXiv.txt": {
        "abstract": "We present a model for determining the dimensionless spin parameter and mass of the black hole remnant of black hole-neutron star mergers with parallel orbital angular momentum and initial black hole spin. This approach is based on the Buonanno, Kidder, and Lehner method for binary black holes, and it is successfully tested against the results of numerical-relativity simulations: the dimensionless spin parameter is predicted with absolute error $\\lesssim 0.02$, whereas the relative error on the final mass is $\\lesssim 2$\\%, its distribution in the tests being pronouncedly peaked at $1$\\%. Our approach and the fit to the torus remnant mass reported in~\\cite{Foucart2012} thus constitute an easy-to-use analytical model that accurately describes the remnant of black hole-neutron star mergers. The space of parameters consisting of the binary mass ratio, the initial black hole spin, and the neutron star mass and equation of state is investigated. We provide indirect support to the cosmic censorship conjecture for black hole remnants of black hole-neutron star mergers. We show that the presence of a neutron star affects the quasinormal mode frequency of the black hole remnant, thus suggesting that the ringdown epoch of the gravitational wave signal may virtually be used to (1) distinguish black hole-black hole from black hole-neutron star mergers and to (2) constrain the neutron star equation of state. ",
        "introduction": "Once a black hole-neutron star (BH-NS) binary is formed, gravitational radiation reaction gradually reduces its orbital separation until the two companions merge and leave behind a remnant consisting of a black hole and, possibly, a hot, massive accretion torus surrounding it~\\cite{ShibataTaniguchilrr-2011-6}. BH-NS binaries have not been observed yet; population synthesis studies, however, suggest that the coalescence of BH-NS systems is likely to occur frequently in the Hubble volume, thus making theoretical studies on the evolution and final state of BH-NS mergers relevant~\\cite{Kalogera2007, Belczynski07, Belczynski08, Oshaughnessy2008, Abadie:2010}. Interest in these systems arises from the fact that they are among the most promising sources for gravitational wave (GW) detectors --- such as LIGO~\\cite{ligowebpage}, Virgo~\\cite{virgowebpage}, KAGRA~\\cite{kagrawebpage}, and the Einstein Telescope~\\cite{Punturo:2010} --- and that they are promising candidates as progenitors of (a fraction of) short-hard gamma-ray bursts~\\cite{Nakar:2007yr, Berger2011}. Further, as NSs in these systems undergo strong tidal deformations, observing GW and/or electromagnetic signals emitted by BH-NS binaries could help shed light on the equation of state (EOS) of matter at supranuclear densities, which is currently unknown~\\cite{Vallisneri00, Ferrari:2009bw, Pannarale2011, Lackey2012}. Finally, comprehending the fate of the material possibly ejected by BH-NS binaries after the NS tidal disruption is relevant in interpreting the observed abundances of the heavy elements that are formed by rapid neutron capture in $r$-processes~\\cite{Lattimer74}. These outflows may additionally be observable due to the radioactive decays triggered by the formation of heavy isotopes, i.e.~``kilonovas,'' or due to the shock they would generate when hitting interstellar medium of sufficiently high density~\\cite{Metzger2012}.  To achieve a full understanding of BH-NS merger events and their physics, numerical-relativity simulations are required. These will ultimately have to include adequate and accurate treatments of general relativity, relativistic (magneto)hydrodynamics, the microphysical EOS, NS crust physics, thermal effects, and nuclear physics reactions. Numerical quasiequilibrium studies~\\cite{Taniguchi05, Grandclement06, Faber06a, Faber05, Tsokaros2007, Taniguchi07, Taniguchi:2008a} and dynamical simulations~\\cite{Janka99a,Ruffert99b, Kluzniak99c, Rosswog05, Loeffler06a, Sopuerta:2006bw, Shibata06d, ShibataUryu:2007, Etienne2007b, ShibataTaniguchi2008, Rantsiou2008, Duez:2008rb, Etienne:2008re, Shibata:2009cn, Duez09, Kyutoku2010, Chawla:2010sw, Foucart2010, Kyutoku2011, Kyutoku2011err, Foucart2011, Lackey2012, Etienne2012, Shibata2012err, Etienne2012b, Deaton2013, Foucart2013a, Kyutoku2013, Foucart2013b, Lovelace2013, Paschalidis2013b} of mixed binary mergers made considerable progress in the last few years. Despite the fact that simulating BH-NS mergers is now possible, these simulations remain nevertheless both challenging and computationally intensive. This has motivated the parallel development of pseudo-Newtonian BH-NS calculations, e.g.~\\cite{Ruffert2010}, and analytical approaches focusing on specific physical aspects of the problem, e.g.~\\cite{Shibata96, Wiggins00, Vallisneri00, Berti08, Hanna2008, Ferrari09, Ferrari:2009bw, Damour:2009wj, Pannarale2010, Pannarale2011, Foucart2012}. Studies of these kinds benefit from their low computational costs which allow them to shed light on questions that cannot be currently addressed with numerical simulations and to provide insight on what happens when the large space of parameters of BH-NS binaries is spanned. They may, in turn, aid numerical relativity by suggesting cases that are particularly interesting to simulate, and by providing information to exploit within the simulations themselves.  In this paper we focus on predicting the spin parameter and mass of the BH remnant of BH-NS coalescing binaries by using a semianalytical approach. While this problem has a fairly long history in the case of coalescing binary black holes~\\cite{Buonanno00a, Damour:2001tu, Hughes:2002ei, Buonanno:06cd, Campanelli:2006gf, Campanelli:2006uy, Campanelli:2006vp, DamourNagar:07a, Tichy:2007gso, Buonanno:07b, Rezzolla-etal-2007, BoyleKesdenNissanke:07, Kesden:2008, Marronetti:2007wz, Rezzolla-etal-2007b, Rezzolla-etal-2007c, Gergely:08, Rezzolla:2008sd, Barausse:2009uz, Kesden2010, Barausse2012b}, no attempt beyond numerical-relativity simulations has yet been made to tackle it in the case of BH-NS mergers. The approach we present and discuss is based on the work of Buonanno, Kidder, and Lehner (BKL) on estimating the final BH spin of a coalescing binary BH with arbitrary initial masses and spins~\\cite{Buonanno:07b}. We choose this simple, phenomenological model as a starting point because it provides good physical insight, and because it is straightforward to modify and extend. Our method may indeed be seen as a generalization of the BKL model to the case in which the lower mass BH is replaced with a NS. For the time being, however, it is restricted to systems in which the BH spin direction is parallel to the orbital angular momentum direction. The closed expression we determine for the final spin parameter automatically yields an estimate of the mass of the BH remnant by means of a method similar to the starting point of Barausse, Morozova, and Rezzolla's calculations on the mass radiated by binary BHs~\\cite{Barausse2012b}, but with modifications inspired, once again, by~\\cite{Buonanno:07b}. The key equations of our approach are Eqs.\\,(\\ref{eq:model-Mf}), (\\ref{eq:model-imp})-(\\ref{eq:fbridge}) and, despite the mathematical complexity of the mixed binary coalescence problem, our method enables us to reproduce the results of numerical-relativity simulations with reasonable accuracy.  The paper is organized as follows. In Sec.\\,\\ref{sec:BKL} we review the BKL approach for binary BHs. In Sec.\\,\\ref{sec:model} we propose an extension of this method in order to predict the spin parameter and mass of BH remnants of BH-NS mergers --- Eqs.\\,(\\ref{eq:model-Mf}), (\\ref{eq:model-imp})-(\\ref{eq:fbridge}) --- and successfully test it against available numerical-relativity data. In Sec.\\,\\ref{sec:results} we gather the results obtained by systematically varying the binary mass ratio, the initial BH spin parameter, and the NS mass and EOS. First, we provide indirect support to the cosmic censorship conjecture and suggest particularly interesting cases to explore with numerical simulations in this context (Sec.\\,\\ref{sec:maxaf}). Then, we show that the NS EOS may leave an imprint on the BH remnant in terms of its final spin and mass (Sec.\\,\\ref{sec:QNMs}). This suggests the idea of inferring the presence of the NS and of constraining its EOS from the ringdown of the BH remnant. Finally, in Sec.\\,\\ref{sec:conclusions}, we draw our conclusions and collect our remarks. ",
        "conclusions": "\\label{sec:conclusions} In this paper we presented a model for predicting the final spin parameter, $a_\\text{f}$, and mass, $M_\\text{f}$, of the BH remnant of BH-NS coalescing binaries in quasicircular orbits, with initial BH spin of arbitrary magnitude and parallel to the orbital angular momentum, with arbitrary mass ratio, and with arbitrary NS mass and cold, barotropic equation of state. The parameter space just outlined could in principle be investigated entirely with numerical-relativity simulations; in practice, however, the process would be very time and resource consuming, because simulations are still very expensive in terms of computational costs.  Our starting point was the phenomenological model of Buonanno, Kidder, and Lehner for the final spin of binary BH mergers~\\cite{Buonanno:07b}, which we modified to account for (1) energy loss via gravitational wave emission during the inspiral and (2) the possible formation of an accretion torus in the case of disruptive mergers. We tested our model by comparing its predictions to the recent numerical-relativity simulation results available in the literature. We were able to achieve good agreement down to a mass ratio of $M_\\text{BH}/M_\\text{NS}=2$, albeit introducing an additional ingredient in the formulation of the model for $2<M_\\text{BH}/M_\\text{NS}<4$ which is currently poorly constrained. We obtained an absolute error on $a_\\text{f}$ of $0.02$, which is compatible with the one of the BKL approach~\\cite{Buonanno:07b,Barausse:2009uz}. For the final gravitational and irreducible (normalized) masses of the BH remnant, $M_\\text{f}$ and $M_\\text{irr,f}$, we found a relative error of $1$\\%. These errors then propagate in the calculation of the $l=m=2$, $n=0$ quasinormal mode frequency $f_{220}^\\text{QNM}$ and damping time $\\tau_{220}^\\text{QNM}$ of the remnant BH, and they yield maximum relative errors of $4$\\% and $5$\\%, respectively. These relative errors are, however, safely $\\leq 2$\\% in the vast majority of test cases. Combining this method with input for the torus remnant mass from the two-parameter model, fitted to existing numerical results, recently reported by Foucart in~\\cite{Foucart2012}, the error $\\Delta a_\\text{f}\\simeq 0.02$ is preserved and so is the behavior of the relative errors on $M_\\text{f}$, $M_\\text{irr,f}$, $f_{220}^\\text{QNM}$, and $\\tau_{220}^\\text{QNM}$ (see Figs.\\,\\ref{FIG:AbsErrorsToy} and \\ref{FIG:HistoMQNMErrors}).  The tests we performed against the available numerical-relativity results were successful, especially considering the limitations of our simple approach. This implies that the outcome of the complicated merger dynamics of BH-NS binaries may be understood in fairly simple terms, at least when the BH initial spin and orbital angular momentum directions are parallel and the inspiral orbit is quasicircular. Equations (\\ref{eq:model-Mf}), (\\ref{eq:model-imp})-(\\ref{eq:fbridge}) presented here along with the fit of Foucart to the torus remnant mass~\\cite{Foucart2012} constitute an easy-to-use analytical model that describes the remnant of BH-NS mergers. Notwithstanding the good performance of the model in the tests, its predictions should be taken with ``a grain of salt,'' as large portions of the parameter space of BH-NS binaries are currently unexplored, hampering a thorough test of our approach.  The approach presented and tested in the first part of this work was then employed to span the space of parameters consisting of the binary (symmetric) mass ratio, the BH initial spin, the NS mass, and the NS equation of state. This enabled us to gain a sense of the behavior of the properties of the BH remnant in this fourfold space of parameters and to pinpoint some interesting aspects which, we believe, deserve being verified and studied in conclusive, quantitative terms with the tools of numerical relativity. The following is a summary of our main results: \\ben[~~~~(i)] \\item We obtained a maximum final spin parameter equal to $1.00$ for the WFF1, APR2, and PS nuclear equations of state, when using the mass ratio $M_\\text{BH}/M_\\text{NS}=10$. $M_\\text{BH}/M_\\text{NS}=2$ yields instead a maximum final spin parameter of $0.99$. Given their absolute error of $0.02$, these predictions are compatible with the $0.98$ maximum found in~\\cite{Kyutoku2011} and provide indirect support to the cosmic censorship conjecture~\\cite{Penrose79}. \\item We discussed the dependence of $a_\\text{f}$ and $M_\\text{f}$ on the NS EOS, claiming that the EOS may leave an imprint on the BH remnant. The quasinormal mode frequency $f_{220}^\\text{QNM}$ of the BH remnant, which depends on $a_\\text{f}$ and $M_\\text{f}$ alone, could thus be used to constrain the NS EOS (Fig.\\,\\ref{FIG:BHBH_QNM}). Deviations from the BH-BH values of $f_{220}^\\text{QNM}$ for symmetric mass ratios $\\gtrsim 0.2$, with maximum deviations between $\\sim 600$Hz and $\\sim 1200$Hz. The excitation of the QNM oscillations does not occur for all mixed binary mergers, but it is likely to appear in the spectrum of the emitted gravitational radiation in the form of a cutoff frequency $f_\\text{cut}\\simeq f_{220}^\\text{QNM}$ for systems with fairly compact NSs, i.e.~for soft EOSs~\\cite{Kyutoku2011}. The possibility of constraining the EOS by measuring $f_\\text{cut}\\simeq f_{220}^\\text{QNM}$ seems therefore complementary to other ideas for posing EOS constraints by means of GW detection, in that these favor constraints on stiff EOSs~\\cite{Vallisneri00, Hinderer09, Pannarale2011, Kyutoku2011}. High-frequency gravitational waves from coalescing binaries may thus turn out to be, once more, very promising in terms of the NS EOS~\\cite{Bauswein2012}. \\een  Future applications of the approach presented in this paper may be to exploit the predicted values of $a_\\text{f}$ and $M_\\text{f}$ to (1) provide the QNM frequencies to be used in the construction of hybrid waveforms for BH-NS systems~\\cite{Kyutoku2011,Lackey2012}, (2) develop phenomenological waveforms for BH-NS systems, (3) study time-frequency characteristics of the emitted radiation~\\cite{Hanna2008}, and (4) build backgrounds for perturbative approaches to the study of the postmerger epoch. Two main extensions are, instead, foreseeable and consist in allowing for a more general initial state for the binary. One may investigate dropping the assumption that (1) the neutron star is initially irrotational and that (2) the initial spin angular momentum of the black hole and the orbital angular momentum are parallel. In the former case, one would have to add a spin angular momentum contribution from the neutron star to the numerator of Eq.\\,(\\ref{eq:model-imp}) and allow for a fraction of this angular momentum to possibly be dissipated prior to the disruption/plunge of the star. A major obstacle, however, would be that, since there are no numerical-relativity simulations with a nonirrotational neutron star initial state, we lack a model to predict $M_\\text{b,torus}$ when the neutron star is initially spinning. As this quantity appears in Eq.\\,(\\ref{eq:model-imp}), being able to predict it would be a fundamental gap to fill in, prior to extending the model discussed in this paper. Regarding cases with nonparallel black hole spin angular momentum and orbital angular momentum, the ISCO-related expressions appearing in Eq.\\,(\\ref{eq:model-imp}) would have to be generalized. Input in this direction has recently started to emerge from the numerical-relativity community via the first simulations of tilted BH-NS mergers, e.g.~\\cite{Foucart2010}, that would serve as test cases. In general, and independently of extending the model, it is important to keep testing the model as more numerical simulations are published in the literature, possibly constraining the ansatz of Eq.\\,(\\ref{eq:fbridge}) and improving the model of~\\cite{Foucart2012}, on which the entire approach relies.  In concluding this work, we would like to stimulate the BH-NS numerical-relativity community to continue investigating different parameter configurations, as this would allow us to better constrain the current version of our model [see the discussion following Eq.\\,(\\ref{eq:model-imp})]. Investigations on more generic initial spin configurations would also be helpful~\\cite{Foucart2010}, as they would allow us to look into extending our approach."
    },
    "1208/1208.1167_arXiv.txt": {
        "abstract": "A survey of radio recombination lines in the Galactic plane with longitude $-32^o < l < +80^o$ and latitude $b<\\pm3^o$ using Ooty Radio Telescope(ORT) at 328 MHz has been reported. ORT observations were made using a New Digital Backend(NDB) augmented to it recently. With NDB ORT had a beam of $2^o.3 \\times 2^o.2 sec(\\delta)$ and a passband of $\\sim$1 MHz in the spectral line mode. The above mentioned Galactic region was divided into $\\sim 2^o \\times 2^o$ patches with the ORT beam pointed to the center. The ORT observations form a study of distribution of extended low-density warm-ionized medium(ELDWIM) in the inner Galaxy using H271$\\alpha$ RL. By obtaining kinematical distances using $V_{LSR}$ of the H271$\\alpha$ RLs the distribution of ELDWIM clouds within the inner Galaxy has been deduced for the region given above. ",
        "introduction": "The preliminary unsuccessful attempts to detect RRL from hydrogen were made by Egorova and Ryzkov (1960) with the Pulkovo telescope.  Successful detection of RRL was made in April 1964 using  an improved radio meter with the 22-m radio telescope at Lebedev Physical institute in Puschino. Sorochenko and Borodzich detected the hydrogen RRL H90$\\alpha$ ($\\lambda$=3.38cm) towards the omega nebula. Independently at about the same time another group(Dravskikh et al 1965) also detected a convincing RRL H104$\\alpha$. Following this many researchers(Lilley et al 1966; Goldberg \\& Dupree 1967; Gottesman \\& Gordon 1970; ) detected RLs. Subsequent attempts(Batty 1976) to detect RRL at lower frequencies($<$ 500MHz) seem (Anantharamaiah 1985)to have failed due to non availiability/development of radio telescope hardware. It was already known(Shaver 1975) in the community that stimulated emission($kT/h\\nu \\times$ spontaneous) would boost RRL at lower frequencies. Apparently RRL at frequencies below 500 MHz were first observed from the Galactic plane(Anantharamaiah 1985) using Ooty Radio Telescope(ORT). A selective survey(Anantharamaiah 1985) of RRL was made in the Galactic plane($b \\sim 0^o$) with longitude $<60^o$ using ORT at 325 MHz(H272$\\alpha$) towards 53 directions. This observation gave numerous RL detections and paved way to a new survey (Roshi \\& Anantharamaiah 2000; Roshi \\& Anantharamaiah 2001)covering the Galactic plane systematically. However these positions have been mostly in the Galactic plane within a latitude of $b<\\pm1^o$. With the introduction of a New Digital Backend(NDB) for the ORT in the recent years it has been possible to carry forward this survey to higher latitudes within less time and better signal to noise ratio(S/N) compared to previous observations. Previous observations have used either the entire telescope with a beam size of $\\sim 2^o \\times 6'$(Anantharamaiah 1985; Anantharamaiah 1986; Roshi \\& Anantharamaiah 2001a) or a couple of modules with a beam size of $\\sim 2^o \\times 2^o$(Roshi \\& Anantharamaiah 2000). In the present observations the NDB could process signals from all the 22 modules(Section 2) of the ORT thus reducing the integration time required for detection. On the other hand  earlier observations have had a larger bandwidth(BW) or more number of  RRL-transitions(Roshi \\& Anantharamaiah 2000) compared to only one observable transition (H271$\\alpha$) with the recent NDB. \\\\  ORT observations were made with a New Digital Backend(NDB) augumented to it recently(Prabu 2010). In the spectral line mode, with NDB, ORT provided a passband of 1.2 MHz and a beam size of $2^o.3 \\times 2^o.2 sec(\\delta)$. The Galactic region between $-32^o < l < +80^o$ and $b<\\pm3^o$ was divided into patches of $2^o \\times 2^o$ with ORT pointed towards the center of each patch. ORT observed these 165 positions distributed in 3 rows and 55 columns for $\\sim$3 hrs per pointing. Some of the positions were skipped due to existence of earlier observations, shortage of telescope time and severe interference at higher declinations. Due to technical reasons H271$\\alpha$ seemed to be the only appropriate RL for ORT with NDB.  Kinematical distances towards H271$\\alpha$ line emitting regions were obtained using a differential Galactic rotation curve(Sofue et al. 2009) which gave a distribution of ELDWIM clouds in the inner Galaxy. ORT observations aimed at obtaining the distribution of ELDWIM in the inner Galaxy and to obtain new RL detections from above and below the Galactic plane at 328 MHz.  \\\\  \\section[]{Observations} Ooty Radio Telescope(ORT) is situated near the town Ooty, south India at a longitude of $283^{o}.3$ and latitude of $11^{o}.4$. ORT(Swarup et al. 1971) is an off-axis parabolic cylinder with a length of 530m and width of 30m. The telescope is located on a hill which has a natural slope of $11^{o}.4$ equal to the geographical latitude of the place. This gives it the feature of equatorial mount. The operating frequency of the telescope is centered at 326.5 MHz with a maximum BW of 15MHz at the front-end. The reflecting surface of the clylinder is made of 1100 stainless steel wires running parallel to each other along the entire length of the telescope. An array of 1056 half-wave dipoles in front of a $90^{o}$ corner reflector forms the primary feed of the telescope. The 1056 dipoles are in groups of 48. The signals recieved by these groups are added in phase to form 22 group outputs, each known as a module. The telescope is divided into northern part and southern part. The northern modules are designated as N1 to N11 and the southern modules as S1 to S11. The beam width due to each module is $2^o.3$ in east-west and $2^o.2 sec(\\delta)$ in the north-south, where $\\delta$ is the declination. This forms the observing mode and beam for the current project of RL observations.\\\\  The RRL observations were made using a new digital backend(Prabu 2010) which could be operated in narrow band mode or broad band(10MHz) mode. The narrow band mode was the spectral line mode which provided a BW of 1.2 MHz. This small BW restricted the observation of only 1 RL at a time. The transistion selected was $H271\\alpha$ which has a rest frequency of 328.5958MHz. ORT has a large front end BW and the NDB's 1.2 MHz BW could have accomodated any of the near by RLs. But due to non availability of a broad band amplifier for the local oscillator(LO) use of ORT's dedicated amplifier with a -3 dB gain BW of 2MHz was employed. This restricted the freedom of deviation from the ORT's central LO feed frequency of 296.5 MHz. A symbolic block diagram of the instrument set up is shown in Figure 1. The observations were carried out using dual frequency switching with a shift of $\\Delta\\nu_{s}$=300 or 400kHz. This magnitude of shift ensured that there was no overlap of the associated carbon RL C271$\\alpha$ with H271$\\alpha$ between 2 shifted spectra. The $V_{LSR}$ difference between C271$\\alpha$ and H271$\\alpha$ is $\\sim$150 km$s^{-1}$. With this arrangement a spectral BW of approximately 300 km$s^{-1}$ in $V_{LSR}$ could be covered. At a frequency of 328 MHz differential frequency roughly transforms into differential velocity according to $\\Delta V = c\\cdot\\Delta\\nu/\\nu$. \\\\  The exact BW of the spectral line mode was decided by the sampling frequency of NDB which was nearly 2.45 MHz, giving a BW of half of sampling frequency as decided by Nyquist sampling rate, 1.225MHz. The resolution of the observed band would then depend upon the number of FFT points performed on the data,  \\begin{equation} \\Delta \\nu_{res} = \\frac{BW}{n_{FFT}}. \\end{equation}  In the present case $n_{FFT}$ = 256 throughout the observations, so $\\Delta \\nu_{res}$ = 4.785 kHz or $\\Delta V_{res}$ = 4.37 km s$^{-1}$. This resolution is acceptable for hydrogen RL which in the present observations was of primary interest. However the same cannot be said for carbon RL. Being heavy and considering its origin from cold regions its line width could be completely contained within this resolution. So the carbon RL is considerably smoothed out.\\\\   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=120mm,height=54mm,angle=0]{ORTdas.eps} \\caption{Symbolic block diagram of ORT instrument setup for H271$\\alpha$ RL observation.} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "1208/1208.4097_arXiv.txt": {
        "abstract": "We search the $\\lambda=1.1$ mm Bolocam Galactic Plane Survey for clumps containing sufficient mass to form $\\sim10^4~\\msun$ star clusters. \\ncandidates\\ candidate massive proto-clusters  are identified in the first Galactic quadrant outside of the central kiloparsec.  This sample is complete to clumps with mass M$_{\\rm clump}>\\mmin$ and radius $r\\lesssim2.5$ pc.  The overall Galactic massive cluster formation rate is $CFR({\\rm M}_{\\rm cluster}>10^4) \\lesssim \\CFR\\  \\permyr$, which is in agreement with the rates inferred from Galactic open clusters and M31 massive clusters.  We find that all massive proto-clusters in the first quadrant are actively forming massive stars and place an upper limit of $\\tau_{starless}<\\tsuplim$~Myr on the lifetime of the starless phase of massive cluster formation.  If massive clusters go through a starless phase with all of their mass in a single clump, the lifetime of this phase is very short. ",
        "introduction": "The Milky Way contains about 150 Globular clusters (GCs) with masses of $10^4$ to over $10^6$ \\msun\\ and tens of thousands of open clusters containing from 100 to over $10^4$ stars.  However, young massive clusters containing $\\gtrsim10^4~\\msun$ of stars are rare, with only a handful known \\citep{PortegiesZwart2010}. While no GCs have formed in the Milky Way within the last 5 Gyr, open clusters that survive many crossing times continue to form.   A few of these clusters have stellar masses greater than $10^4$ M$_{\\odot}$ and therefore qualify as young massive clusters \\citep[YMCs;][]{PortegiesZwart2010}.   YMCs must either form from clumps having masses greater than and sizes comparable to the final cluster  or be formed from a larger, more diffuse reservoir, in which case massive protocluster clumps may be rare or nonexistent  \\citep{Kennicutt2012}.    Massive proto-clusters (MPCs) are massive clusters (M$_{\\rm cluster}>10^4$ \\msun) in the process of forming from a dense gas cloud.  In \\citet{Bressert2012}, we examine the theoretical properties of MPCs: MPCs are assumed to form from massive, cold starless clumps analagous to pre-stellar cores \\citep{Williams2000}.  In this paper, we refer to two classes of objects: starless MPCs, which have very low luminosity and do not contain OB stars, and MPCs, which are gas-rich but have already formed OB stars.  The only currently known starless MPC is G0.253+0.016, which lies within the dense central molecular zone and is subject to greater environmental stresses than similar objects in the Galactic plane \\citep{Longmore2012}.  Because massive clusters contain many massive stars, at some point during their evolution ionization pressure will prevail over protostellar outflows as the dominant feedback mechanism.  Other sources of feedback are less than ionization pressure up until the first supernova explosion \\citep{Bressert2012}.  These proto-clusters must have masses M$_{\\rm clump}>{\\rm M}_{*}/SFE$ \\footnote{We define a star formation efficiency $SFE={\\rm M}_{\\rm *,final} / {\\rm M}_{\\rm gas,initial}$.}, or about $3\\ee{4}$ \\msun\\ for an assumed SFE=30\\% (an upper limit on the star formation efficiency), confined in a radius $r\\lesssim2.5$ pc, in order to remain bound against ionization feedback.  These properties motivate our search for proto-clusters in the Bolocam Galactic Plane Survey \\citep[BGPS;][ \\url{http://irsa.ipac.caltech.edu/data/BOLOCAM_GPS/}]{Aguirre2011}.   The distinction between relatively short-lived `open clusters' and long-lived ($t\\gtrsim1$ Gyr) bound clusters occurs at about $10^4$ \\msun \\citep{PortegiesZwart2010}.  Clusters with ${\\rm M}_{\\rm cluster} < 1\\ee{4} \\msun$~will be destroyed by interactions with giant molecular clouds over the course of a few hundred million years after they have dispersed their gas \\citep{Kruijssen2011}, while clusters with ${\\rm M}_{\\rm cluster}\\gtrsim10^4 \\msun$ may survive $\\gtrsim 1$ Gyr.  Closer to the Galactic center, within approximately a kiloparsec, all clusters will be destroyed on shorter timescales by strong tidal forces or interactions with molecular clouds.  In the Galaxy, there are few known massive clusters. \\citet{PortegiesZwart2010} catalogs a few of them, of which NGC 3603, Trumpler 14, and Westerlund 1 and 2 are the likely descendants of the objects we investigate.  These clusters have $r_{eff} \\lesssim 1$ pc, $M\\sim10^4$ \\msun, and ages $t\\lesssim4$ Myr.  We present a census of their ancestral analogs. ",
        "conclusions": "\\label{sec:conclusions}  Using the BGPS, we have performed the first flux-limited census of massive proto-cluster candidates.  We found \\ncandidates\\ candidates that will be part of the next generation of open clusters and \\nMPC\\ that could form massive clusters similar to NGC 3603 (${\\rm M}_{\\rm cluster} > 10^4$ \\msun).   We have measured a Galactic massive cluster formation rate $CFR({\\rm M}_{\\rm cluster}>10^4) \\lesssim \\CFR\\  \\permyr$\\ assuming that clusters are equally likely to form everywhere within the range 1 kpc $ < R_{gal} < $ 15   kpc. The observed MPC counts are consistent with observed cluster counts in Andromeda scaled up by $SFR_{M31} / SFR_{MW}$ assuming a formation timescale of 2 Myr.  Despite this survey being the first sensitive to pre-star-forming MPC clumps, none were detected.  This lack of detected pre-star-forming MPCs suggests a timescale upper limit of about $\\tau_{starless}<\\tsuplim$ Myr for the pre-massive-star phase of massive cluster formation, and hints that massive clusters may never form highly condensed clumps ($\\bar{n}\\gtrsim10^4~\\percc$) prior to forming massive stars. It leaves open the possibility that massive clusters form from large-scale ($\\gtrsim 10$ pc) accretion onto smaller clumps over a prolonged ($\\tau > 2$ Myr) star formation timescale.   Observations are needed to distinguish competing models for MC formation: Birth from isolated massive proto-cluster clumps, either compact and rapid or diffuse and slow, or from smaller clumps that never have a mass as large as the final cluster mass. This sample of the \\ncandidates\\ most massive proto-cluster clumps in the first quadrant (where they can be observed by both the VLA and ALMA) presents an ideal starting point for these observations."
    },
    "1208/1208.1351_arXiv.txt": {
        "abstract": "We present numerical simulations of the kinematic induction equation in order to examine the dynamo efficiency of an axisymmetric von K{\\'a}rm{\\'a}n--like flow subject to time-dependent nonaxisymmetric velocity perturbations.  The numerical model is based on the setup of the French von K{\\'a}rm{\\'a}n-sodium dynamo (VKS) and on the flow measurements from a water experiment conducted at the University of Navarra in Pamplona, Spain. The principal experimental observations that are modeled in our simulations are nonaxisymmetric vortexlike structures which perform an azimuthal drift motion in the equatorial plane. Our simulations show that the interactions of these periodic flow perturbations with the fundamental drift of the magnetic eigenmode (including the special case of nondrifting fields) essentially determine the temporal behavior of the dynamo state.  We find two distinct regimes of dynamo action that depend on the (prescribed) drift frequency of an ($m=2$) vortexlike flow perturbation. For comparatively slowly drifting vortices we observe a narrow window with enhanced growth rates and a drift of the magnetic eigenmode that is synchronized with the perturbation drift. The resonance-like enhancement of the growth rates takes place when the vortex drift frequency roughly equals the drift frequency of the magnetic eigenmode in the unperturbed system. Outside of this small window, the field generation is hampered compared to the unperturbed case, and the field amplitude of the magnetic eigenmode is modulated with approximately twice the vortex drift frequency. The abrupt transition between the resonant regime and the modulated regime is identified as a spectral exceptional point where eigenvalues (growth-rates and frequencies) and eigenfunctions of two previously independent modes collapse.  In the actual configuration the drift frequencies of the velocity perturbations that are observed in the water experiment are much larger than the fundamental drift frequency of the magnetic eigenmode that is obtained from our numerical simulations. Hence, we conclude that the fulfillment of the resonance condition might be unlikely in present day dynamo experiments. However, a possibility to increase the dynamo efficiency in the VKS experiment  might be realized by an application of holes or fingers on the outer boundary in the equatorial plane. These mechanical distortions provoke an anchorage of the vortices at fixed positions thus allowing an adjustment of the temporal behavior of the nonaxisymmetric flow perturbations. ",
        "introduction": "Cosmic magnetic fields are ubiquitous phenomena that are intrinsically coupled to most astrophysical objects like planets, stars, or galaxies. The origin of these fields involves the formation of electrical currents by means of a complex flow of a conducting fluid or plasma. This process, the so-called dynamo effect, is necessarily three dimensional and nonlinear, which makes an analytical or numerical approach difficult.  Meanwhile, fluid-flow-driven generation of magnetic fields has also been obtained in laboratory experiments providing a complementary tool to astronomical observations or numerical simulations.  However, whereas astrophysical dynamo action is comparably easy because of the large dimensions of the involved flows, its experimental realization requires considerable technical efforts \\cite{ISI:000262272000001}.  An important obstacle for the occurrence of laboratory dynamo action arises from the scaling behavior of the power that is required to drive a flow with a requested magnetic Reynolds number, ${\\rm{Rm}}$.  For turbulent flows this power scales $\\propto {\\rm{Rm}}^3$ so a reduction of the critical ${\\rm{Rm}}$ for the onset of dynamo action is most important to achieve magnetic self-excitation at all.    So far, dynamo experiments based on a flow of a conducting fluid have been successfully conducted in Riga \\citep{2000PhRvL..84.4365G}, Karlsruhe \\citep{abs}, and Cadarache \\citep{2007PhRvL..98d4502M}.  The first two facilities made use of a more or less predetermined fluid flow essentially fixed by the forcing and the shape of the internal tubes.  Note, however, that, at least in the Riga dynamo experiment, the saturation process involved a nontrivial back-reaction effect of the magnetic field that changes the geometry of the flow. Such effects might be even more pronounced in the Cadarache von-K{\\'a}rm{\\'a}n-sodium (VKS) dynamo. In that experiment, the flow driving by two opposing impellers provides more freedom for the development of a saturated turbulent state, in which the back-reaction of the magnetic field on the fluid can strongly modify the geometry and dynamics of the flow.  In an idealizing model the mean axisymmetric flow between counter-rotating impellers comprises two toroidal and two poloidal eddies (so-called {\\it{s2t2}} topology), and it is well known that this flow is able to drive a dynamo \\cite{1989RSPSA.425..407D,1998PhRvE..58.7397O}.  Various attempts in different geometries have been made (numerically as well as experimentally) in order to examine dynamo action driven by such a flow \\citep{forest2002,2003EPJB...33..469M,2004mag...bourg,2005physics..11149S, 2005PhFl...17k7104R,2007EL.....7759001B,2007PhRvL..98d4502M,2008EL.....8229001G, 2008PhRvL.101n4502G,2009NJPh...11a3027R,2009PhRvE..80e6304R,2009PhFl...21c5108M, 2010GApFD.104..207R,2011PhPl...18c2110K,2011PhRvL.106y4502K}. However, so far, experimental dynamo action driven by a von K{\\'a}rm{\\'a}n--like flow is obtained only at the VKS facility and only when at least one of the flow-driving impellers is made of soft iron with a large relative permeability.  Kinematic simulations of the Cadarache dynamo indicate a close linkage between the exclusive occurrence of dynamo action in the presence of soft iron impellers and the observed axisymmetry of the magnetic field \\cite{2010PhRvL.104d4503G,2012NJPh...14e3005G}.  Nevertheless, a fully satisfactory explanation of the working principle of this dynamo is still missing and it is still unclear whether the present experiments will ever be able to achieve growing equatorial dipole modes, which constitute the magnetic field geometry that has been expected from kinematic simulations with an axisymmetric flow field.      An improvement of present numerical models may require the explicit consideration of coherent nonaxisymmetric structures that repeatedly have been observed in water experiments using a von K\\'arm\\'an--like flow driving \\cite{2007PhRvL..99e4101D, 2008JFM...601..339R, 2009PhFl...21b5104C}.  Nonaxisymmetric time-periodic flows with a dominant azimuthal wave number $m=2$ have also been found in 3D simulations of {\\it{s2t2}} flows in spherical geometry \\cite{2009PhRvE..80e6304R}.  Kinematic dynamo simulations using various manifestations of these velocity fields showed a surprising diversity of behavior patterns, however, self-generation of magnetic energy was found only when the time-dependent flow field was taken into account, whereas the simulations with the time-averaged flow or with different snapshots of the velocity field did not exhibit dynamo action. A similar behavior has also been found previously in an ideal two-dimensional model by the authors of Ref. \\cite{dormy_gerard-varet_2008}, who examined the induction action of a uniform shear flow perturbed by an periodic variation on intermediate time scales. The authors found a perpetual amplification even for very small perturbation amplitudes and concluded that tiny distortions $\\sim \\rm{Rm}^{-1}$ can be sufficient to alter the ability of a flow to provide for dynamo action. This type of dynamo action has been attributed to {\\it{non-normal growth}} in Ref. \\cite{2008PhRvL.100l8501T}, where it was shown that an appropriate mixing of nonorthogonal eigenstates through a time-dependent linear operator can lead to growing modes even if the contributing eigenstates alone correspond to decaying solutions in a stationary system.  The main objective of our study is the behavior of the dynamo efficiency of a cylindrical VKS-like system subject to nonaxisymmetric velocity perturbations with a single azimuthal wavenumber $m=2$.  Such velocity modes were observed in a water experiment conducted at the University of Navarra in Pamplona in order to analyze the influence of slowly evolving large-scale flow on the occurrence of dynamo action \\cite{2007PhRvL..99e4101D,ISI_000271475800018}.  Here, we utilize the essential features of the measured flow field as the basis input for numerical simulations of the kinematic induction equation. Typical input parameters that are systematically varied are the flow amplitude (in terms of the magnetic Reynolds number) and the azimuthal drift motion of the implied nonaxisymmetric velocity perturbation. From the simulation data we extract the leading eigenmodes and the related eigenvalues in terms of growth rates and frequencies that describe field amplitude modulations and/or azimuthal field drift. Interestingly, for comparably low drift frequencies of the velocity perturbation, we, first, observe a phase locking of the magnetic eigenmode drift with the vortex drift which is replaced, for higher perturbation drift frequency, by the appearance of a time-modulated magnetic eigenmode. By analyzing the involved growth rates and frequencies in the phase-locked regime and in the modulated regime, we identify the transition between them as a spectral exceptional point where eigenvalues and eigenfunctions of two modes coincide \\cite{katobook,2004CzJPh..54.1039B}.  The observed behavior is in close analogy with typical (resonant) mechanical systems subject to periodic forcing, like, e.g., spinning disk systems \\cite{1992JAM....59..390C}, or to the behavior observed in the stability study of water waves \\cite{1986RSPSA.406..115M} and we will see, by analyzing the solution of a simple Mathieu equation, that the observed spectral structure is quite generic for systems under the influence of periodic forcing. Comparable effects have also been found in mean-field dynamos of $\\alpha\\omega$-type that were designed to explain the bisymmetric field pattern observed in spiral galaxies.  In these models a periodic perturbation is caused by density waves due to spiraling arms, and a parametric resonance (also called {\\it{swing excitation}}) is observed when the frequency of the perturbing velocity pattern is twice the oscillation frequency of the (axisymmetric) dynamo \\cite{1990MNRAS.244..714C,1992A&A...264..319S, 1996A&A...308..381M, 1999A&A...347..860R, 2002MNRAS.334..925S}.  In the present paper, we show that a facilitation of dynamo action by periodic flow perturbations is also possible in more complex three-dimensional models that include magnetic diffusivity and potentially can be applied to existing dynamo experiments. In contrast to the dynamo models from Refs. \\cite{2009PhRvE..80e6304R} and \\cite{2008PhRvL.100l8501T} the observed increase of the growth rate occurs already without involving time-periodic states, which makes an interpretation in terms of non-normal growth (as in Ref. \\cite{2008PhRvL.100l8501T}) rather implausible. ",
        "conclusions": "We have examined kinematic dynamo action of a von K\\'arm\\'an-like flow of a conducting fluid in a cylindrical container.  When the flow breaks the ideal equatorial symmetry of the system, the critical magnetic Reynolds number for the onset of dynamo action increases with the amount of symmetry breaking and a time dependence is introduced in terms of an azimuthal drift motion of the dominant dynamo eigenmode. The frequency of this drift increases with the amount of symmetry breaking as well as with the magnetic Reynolds number.  \\begin{figure}[t!] \\includegraphics[width=16cm]{./giesecke_nonaxi_vks_pic14.ps} \\caption{(Color online) Alignment of magnetic eigenmode and velocity perturbation for different values of the vortex drift frequency $\\omega_{\\rm{v}}$. All runs stem from the resonant regime. The color-coded (gray shaded) structure denotes $B_\\varphi$ and the contour lines show the axial velocity perturbation $u_z^{\\rm{v}}$. Left: $a=0$; right: $a=0.62$}\\label{fig::phase_fieldmode_velmode} \\end{figure}  The main focus of our examinations has been on the interaction of this field drift with a nonaxisymmetric time-dependent velocity perturbation and the resulting impact on dynamo action.  In summary, what we observe is the following: The temporal behavior of the system is governed by three different time scales: the decay/growth time, the magnetic field drift and the vortex drift, which, in contrast to the first two, represents an imposed quantity.  For rather slow vortex drift frequencies, the first magnetic eigenmode (dominated by an $(m=1)$ component) is enslaved by the drifting vortex pattern, hence, we see here $\\omega_{\\rm{f}}\\sim \\omega_{\\rm{v}}$.  Connected with the linear frequency relationship in this regime we observe a parabolic shape of the growth rate, with a maximum close to the resonance point where the vortex drift frequency would roughly correspond to the eigenfrequency of the magnetic field for the unperturbed, axisymmetric flow, and a quadratic reduction of the growth rate nearby this point. Outside of the resonant regime the field amplitude and the field drift are modulated with twice the frequency of the vortex drift. Phenomenologically, the development of growth rates and frequencies can be described by a Mathieu-like equation.  The observed behavior can also be explained on the basis of simple physical principles. For sufficiently slowly drifting vortices the system adjusts itself to an optimum state and the (azimuthal) phase between magnetic eigenmode and vortex pattern remains fixed so that the field growth becomes maximal.  This state is essential characterized by the alignment between the magnetic eigenmode and the nonaxisymmetric velocity mode which is roughly the same independently of symmetry breaking or vortex drift (see Fig.~\\ref{fig::phase_fieldmode_velmode}).  Increasing $\\omega_{\\rm{v}}$, the magnetic eigenmode cannot follow the ever faster $(m=2)$ vortex drift, but \"bethinks\" of its own eigenvalue (in the unperturbed state) to which it converges in the limit $|\\omega_{\\rm{v}}| \\rightarrow \\infty$.  In doing so, it will be \"beaten\" by the $m=2$ vortex mode with an ever increasing frequency $2\\omega_{\\rm{v}}$, which explains the occurrence of the second frequency involved.    The simulations presented in this study show similarities with the results from Ref. \\cite{2009PhRvE..80e6304R} where dynamo action was examined using a flow field obtained from nonlinear simulations of spherical {\\it{s2t2}} flow in a sphere. Kinematic dynamo simulations using the time averaged flow or different snapshots of the velocity field did not exhibit dynamo action, whereas this was indeed the case when considering the time-dependent flow. The exclusive occurrence of dynamo action with a time-dependent flow was interpreted as dynamo action based on {\\it{non-normal growth}} originally described in Refs. \\cite{dormy_gerard-varet_2008,2008PhRvL.100l8501T}. Looking at the details, a number of differences become visible between the study of Ref. \\cite{2009PhRvE..80e6304R} and the model presented here. The key result of Ref. \\cite{2009PhRvE..80e6304R} is a magnetic field amplification induced by two counterpropagating $m=2$ waves that were anchored at the poles of a sphere. However, in their study, the authors did not find any clear sharp resonant-like regime. Furthermore, they only found a more complex temporal (i.e. cyclic) behavior in one special case with a coupled Navier-Stokes and induction equation which leads to a phase jump in the alignment between the nonaxisymmetric velocity mode and the magnetic eigenmode enforced by the back-reaction of the Lorentz force on the velocity field. In all other runs presented in Ref. \\cite{2009PhRvE..80e6304R} the orientation of the magnetic eigenmode remains fixed in space and time and no azimuthal drift is observed (this behavior is similar to the results presented in Ref. \\cite{2008PhRvL.100l8501T} where only real eigenvalues of the leading eigenmodes were obtained). The most probable explanation for the distinct behavior results from the differences in the hydrodynamic base of the model of Ref. \\cite{2009PhRvE..80e6304R}, which does not show any equatorial symmetry breaking (thus, the basic magnetic state is stationary and non-drifting). Moreover, the wavelike distortions in Ref. \\cite{2009PhRvE..80e6304R} consist of two counterpropagating $m=2$ patterns that were anchored at the poles of a sphere which most probably inhibits the phase-locking phenomenon observed in our study. As a further difference, we observe a resonant behavior even in runs without equatorial symmetry breaking and with stationary vortices, i.e., in systems without any time dependence so no time-dependent contribution is available that may provide for a mixture of non-normal modes.  Hence, the role of non-normal growth in our models and the comparability with the model of Ref. \\cite{2009PhRvE..80e6304R} must for now remain an open question.     It is difficult to conclude if the resonance effect can be realized in existing dynamo experiments.  Although a coincidence of forcing frequency (the frequency of the drift motion) and magnetic field frequency cannot be ruled out, the resonance condition is a quite particular case and most probably can only be fulfilled by chance. This is particularly true in the case of a more realistic non-linear analysis, in which the Navier-Stokes and the induction equation are coupled so the drift of the vortex pattern might strongly be influenced by back-reaction of the magnetic field.  Furthermore, in the examined VKS-like configuration, the vortex drift frequencies observed in the water experiment are far away from the resonance condition (at least for reasonable values of the equatorial symmetry breaking) so it is unlikely that the vortices affect the VKS dynamo in its actual setup.  Nevertheless, it might be suggestive to change the large-scale flow geometry in order to adjust the relation of vortex drift and field drift, e.g., by changing the aspect ratio \\cite{1993Ap&SS.208..245K} or by fixing the vortices using some wall inhomogeneity. From Table~\\ref{tab::freq} it follows immediately that the most promising state to realize the resonance in the experiment would be in a configuration that suppresses the equatorial symmetry breaking (e.g., using a ring in the equatorial plane). In that case, the resonance maximum occurs for nondrifting vortices ($\\omega_{\\rm{v}}=0$), providing a reduction of $\\sim$15\\% (from ${\\rm{Rm}}^{\\rm{c}}=59.7$ to ${\\rm{Rm}}^{\\rm{c}}=50.5$). This configuration is definitely achievable in the experiment where the vortices can be anchored by mounting some inhomogeneity on the outer cylinder wall like, e.g., holes or fingers. However, even when the resonance condition theoretically is adjusted, it remains unclear whether a parametric resonance condition can be fulfilled because in the highly turbulent regime the vortex position undergoes considerable fluctuations. These fluctuations (which indeed are observed in the water experiment) could be modeled by introducing a random phase in the equations for the nonaxisymmetric flow perturbation but, in contrast to amplitude fluctuations, such phase noise is known to prevent the occurrence of a parametric resonance \\cite{petrel_fauve_2006}."
    },
    "1208/1208.4929_arXiv.txt": {
        "abstract": "In order to understand the heating of the solar corona it is crucial to obtain empirical information on the magnetic field in its lower boundary (the transition region). To this end, we need to measure and model the linear polarization produced by scattering processes in strong UV lines, such as the hydrogen Ly$\\alpha$ line. The interpretation of the observed Stokes profiles will require taking into account that the outer solar atmosphere is highly structured and dynamic, and that the height of the transition region may well vary from one place in the atmosphere to another. Here we report on the Ly$\\alpha$ scattering polarization signals we have calculated in a realistic model of an enhanced network region, resulting from a state-of-the-art radiation MHD simulation. This model is characterized by spatially complex variations of the physical quantities at transition region heights. The results of our investigation lead us to emphasize that scattering processes in the upper solar chromosphere should indeed produce measurable linear polarization in Ly$\\alpha$. More importantly, we show that via the Hanle effect the model's magnetic field produces significant changes in the emergent $Q/I$ and $U/I$ profiles. Therefore, we argue that by measuring the polarization signals produced by scattering processes and the Hanle effect in Ly$\\alpha$ and contrasting them with those computed in increasingly realistic atmospheric models, we should be able to decipher the magnetic, thermal and dynamic structure of the upper chromosphere and transition region of the Sun. ",
        "introduction": "}  In order to understand the heating of the solar corona it is crucial to obtain empirical information on the strength and orientation of the magnetic field in its lower boundary (the transition region), where over very short distances (${\\sim}100$ km) the temperature suddenly rises from $<10^4$ K to $>10^6$ K and the atmospheric plasma changes from partially to practically fully ionized \\citep[e.g.,][]{golub-pasachoff10}. To this end, we need to measure and model observables sensitive to the magnetic field of the solar transition region.  In a recent paper, \\citet{jtb-stepan-casini11} argued that the hydrogen Ly$\\alpha$ line is expected to show measurable scattering polarization when observing the solar disk, and that via the Hanle effect the line-center linear polarization amplitude shows a good sensitivity to magnetic field strengths between 10 and 100 G, approximately. These conclusions were reached through detailed radiative transfer calculations in semi-empirical and hydrodynamical  models of the solar atmosphere, assuming complete frequency redistribution (CRD) and neglecting quantum interference between the two upper levels of the Ly$\\alpha$ line (see \\citet{stepanjtb11a} for details on the atomic model and numerical method of solution). As shown by \\citet{belluzzi-trujillobueno-stepan}, these last two approximations are suitable for estimating the polarization signals at the core of Ly$\\alpha$, which is the spectral line region where the Hanle effect in Ly$\\alpha$ operates.\\footnote{Partial frequency redistribution (PRD) and $J$-state interference can, however, produce large linear polarization signals in the wings of the $Q/I$ profile of Ly$\\alpha$, especially for close-to-the-limb line of sights \\citep[see][]{belluzzi-trujillobueno-stepan}.}  The above-mentioned investigations were carried out using one-dimensional models of the extended solar atmosphere, such as the semi-empirical models of \\citet{falc} and the hydrodynamical models of \\citet{carlsson-stein-1997}. However, both observations \\citep[e.g.,][]{vourlidas-2010} and simulations \\citep[e.g.,][]{leenaarts12} show that the outer solar atmosphere is highly structured and dynamic, departing radically from a uniform, plane-parallel configuration. Thus, the height of the transition region must vary from one place in the atmosphere to another, most probably delineating a highly corrugated surface. It is in this type of spatially complex plasmas where strong transition region lines like Ly$\\alpha$ originate. The presence of horizontal atmospheric inhomogeneities in a stellar atmosphere breaks the axial symmetry of the radiation field within the medium, and this type of symmetry breaking may produce changes in the scattering polarization signals that could be confused with those caused by the Hanle effect \\citep[e.g.,][]{manso-jtb11,anusha-nagendra}. It is therefore important to investigate the problem of scattering polarization and the Hanle effect in Ly$\\alpha$ using spatially complex models of the outer solar atmosphere, where the physical quantities may vary abruptly not only along the radial direction but also horizontally. To this end, \\citet{stepanjtb12a,stepanjtb12b} have developed a three-dimensional (3D), multilevel radiative transfer code for doing simulations of the spectral line polarization caused by scattering processes and the Hanle and Zeeman effects within the framework of the quantum theory of polarization described in \\citet{landi-landolfi04}.  The present paper represents a first step towards that goal. Given the complexity of the radiative transfer problem of the Hanle effect in 3D model atmospheres here we start by considering a two-dimensional (2D) model atmosphere characterized by spatially complex variations of the physical quantities at chromospheric and transition region heights, taken from the 3D radiation magneto-hydrodynamic (MHD) simulations mentioned below. Of particular interest here is to compare the $Q/I$ and $U/I$ profiles computed assuming $B=0$ G at each spatial grid point of the considered MHD atmospheric model with the fractional linear polarization signals calculated taking into account the Hanle effect caused by the model's magnetic field. ",
        "conclusions": "}  In order to obtain quantitative information on the magnetic field of the upper chromosphere and transition region of the Sun we need to measure and model the linear polarization caused by scattering processes in some strong resonance lines, such as the hydrogen Ly$\\alpha$ line at 1216\\,\\AA. This is because via the Hanle effect such polarization is sensitive to the strength and orientation of the magnetic field expected for the solar transition region plasma (with $B\\lesssim 100$\\,G). The interpretation of the observed Stokes profiles will require taking into account that the outer solar atmosphere is highly structured and dynamic. For this reason it is important to do numerical simulations of the scattering polarization that is produced in realistic MHD models of the solar atmosphere, paying particular attention to understanding the impact of the model's magnetic field on the emergent $Q/I$ and $U/I$ profiles.  In this Letter we have made a first step towards that goal, by solving the multilevel radiative transfer problem of the scattering polarization and Hanle effect in Ly$\\alpha$ in a 2D snapshot taken from the 3D radiation MHD simulation of an enhanced network region described in Leenaarts et al. (2012). To this end, we applied the 3D radiative transfer code described in \\citet{stepanjtb12a,stepanjtb12b}.  Our model calculations of the line-core $Q/I$ and $U/I$ signals at each point on the model's surface show that the horizontal inhomogeneities of the chosen atmospheric model produce a significant impact on the local scattering polarization amplitudes, and that these linear polarization signals are significantly different in the absence and in the presence of the Hanle effect produced by the model's magnetic field. Moreover, the spatially-averaged $Q/I$ and $U/I$ signals show also an interesting sensitivity to the model's magnetic field. We therefore conclude that by measuring the polarization signals produced by scattering processes and the Hanle effect in Ly$\\alpha$ and contrasting them with those computed in increasingly realistic atmospheric models, we should be able to decipher the magnetic, thermal and dynamic structure of the upper chromosphere and transition region of the Sun."
    },
    "1208/1208.6437_arXiv.txt": {
        "abstract": "The expansion of a  superbubble is investigated both analytically and numerically. Our model implements the thin layer approximation in a vertical profile of density as given by an isothermal self-gravitating disk. A precise  comparison with the results of numerical hydro-dynamics  is given. Analogies  are drawn with the Kompaneets equation  that  includes the  quadratic hyperbolic-secant law in the list  of the plane-parallel stratified media. An  astrophysical application  is made to  the  superbubble connected with the two  worms 46.4+5.5 and  39.7+5.7. The effects of the rotation  of the galaxy on the simulated radius and on the velocity are  introduced. The worms with their strong   limb-brightening visible  on astronomical maps are explained in the framework of image theory. ",
        "introduction": "The superbubble (SB) plays a relevant role in astrophysics because (i) it transports material from the galactic plane  up to a great galactic height, (ii) it can be a site for the acceleration of cosmic rays, see \\cite{Higdon2005,Higdon2006,Butt2008,Ferrand2010}. SBs have been observed in various bands: for  the H~I maps, see \\cite{Oey2002}, for  the optical and HI observations, see \\cite{Oey2002}, and for  the  X-ray  maps, see \\cite{Chu2008,Rodriguez2011,Jaskot2011}; often the image shows a strong   limb-brightening,   which indicates that the emitting layer is thin. The  worm is another observed feature that may, or may not, be associated with a wall of an SB. Galactic worms were first identified as irregular, vertical columns of atomic gas stretching from the galactic plane; now, similar structures are found in radio continuum and infrared maps, see for example~\\cite{Koo1992,English2000,Baek2008}. The models  that  explain  the SB  as being due to the combined  explosions  of supernova in a cluster  of massive stars will now be briefly reviewed. The hydrodynamical approximation, with the inclusion of interstellar density gradients, can  produce  a blowout  into the galactic halo, see  \\cite{MacLow1989,Melioli2009}. Expansion in the presence of magnetic fields has  been implemented in various  magneto-hydrodynamic codes, see \\cite{Tomisaka1992,Rafikov2000}. In  semi-analytical calculations, the thin layer approximation can be the key to obtaining the expansion of the SB: see,  for example, \\cite{McCray1987,McCray1987b,MacLow1988}. The thin layer approximation allows of finding the equation of motion for an expansion in the framework of momentum conservation, see  \\cite{Dyson1997}, when the density of the surrounding medium is constant. The case of an expansion in a medium with variable density is more complex and an exponential and a power law vertical profile have been analysed, see  \\cite{Zaninetti2010h}. The exponential and vertical profiles in density do not correspond to some physical process of equilibrium. The case of an isothermal self-gravitating disk (ISD) is an equilibrium  vertical profile which can be coupled with momentum conservation. The models cited leave some questions unanswered or only partially answered: \\begin {itemize} \\item  Is it possible to calibrate the vertical profile of an ISD? \\item  Is it possible to deduce an analytical formula for the temporal evolution of an SB in the presence of a vertical profile density as given by an ISD? \\item Is it possible to  deduce numerical results for an SB when the  expansion starts at a given galactic height? \\item What is  the influence of galactic rotation on the temporal evolution of an SB? \\item Can we explain the worms as a particular effect using image   theory   applied to SBs? \\end{itemize} In order  to answer these questions, Section \\ref{sec_motion} reviews the standard equation for  momentum conservation in an advancing shell, Section \\ref{sec_asymmetry}  introduces a vertical  profile   in the number of particles as  given by an ISD which  models an aspherical expansion, Section \\ref{sec_applications} applies the new  law  of motion to  the SB associated with GW~46.4+5.5, and Section \\ref{sec_image} contains detailed information on how to build an image of a  SB as  seen from the equatorial plane as well  from the poles. ",
        "conclusions": "{\\bf Law of motion} The temporal evolution  of  an SB  in a medium with constant  density is  characterized  by a spherical symmetry. The presence  of a thin self-gravitating disk of gas which is characterized by a Maxwellian distribution in velocity and  distribution of density which varies only in the z-direction    produces an axial  symmetry in the temporal evolution of an SB. The resulting Eq. (\\ref{fundamental}) has an analytical form which can be solved numerically when $z_{\\mathrm{OB}}$ =0. The case of  $z_{\\mathrm{OB}}\\neq 0$ can be attached solving    two recursive equations, see (\\ref{recursive}). These complex shapes can  also be  modeled by the Kompaneets equation when   a quadratic hyperbolic-secant is adopted. Both  the  approximation here used and the Kompaneets equation can model the shape of the SB without invoking the collision of two wind-blown SBs, see \\cite{Ntormousi2011}. The code developed  here runs in less than a minute on a LINUX -$2.66$GHz processor and can be an alternative to a purely numerical 2D or 3D model.  {\\bf Synchrotron Emission}  Here we assumed that the conversion from the flux of kinetic energy into  non thermal luminosity is regulated by a constant. The exact value of this constant can be deduced once the number density, radius, thickness, and velocity  of the advancing shell are provided by observational astronomy. The values of the fraction of flux deposited in the various  astronomical  bands, $f_c$,  are  given in Table \\ref {tablecolors}.  {\\bf Images}  The emissivity in the thin  advancing  layer is assumed to be  proportional  to the  rate of  kinetic energy, see  Eq.~(\\ref{fluxkineticenergy}), where  the density  is assumed to be proportional  to the swept material. This  assumption explains the strong   limb-brightening visible  on the astronomical maps of SBs and allows associating the observed worms to the non detectable SBs. As an example, if the threshold of the observable flux at a given wavelength  is bigger  than the emitted flux  at the centre of an SB,  we  will  detect only the external regions, the so called worms, see  Fig. \\ref{f12}."
    },
    "1208/1208.4604_arXiv.txt": {
        "abstract": "Molecular hydrogen ($\\H2$) is the primary component of the reservoirs of cold, dense gas that fuel star formation in our galaxy. While the $\\H2$ abundance is ultimately regulated by physical processes operating on small scales in the interstellar medium (ISM), observations have revealed a tight correlation between the ratio of molecular to atomic hydrogen in nearby spiral galaxies and the pressure in the mid-plane of their disks. This empirical relation has been used to predict $\\H2$ abundances in galaxies with potentially very different ISM conditions, such as metal-deficient galaxies at high redshifts. Here, we test the validity of this approach by studying the dependence of the pressure -- $\\H2$ relation on environmental parameters of the ISM. To this end, we follow the formation and destruction of $\\H2$ explicitly in a suite of hydrodynamical simulations of galaxies with different ISM parameters. We find that a pressure -- $\\H2$ relation arises naturally in our simulations for a variety of dust-to-gas ratios or strengths of the interstellar radiation field in the ISM. Fixing the dust-to-gas ratio and the UV radiation field to values measured in the solar neighborhood results in  fair agreement with the relation observed in nearby galaxies with roughly solar metallicity. However, the parameters (slope and normalization) of the pressure -- $\\H2$ relation vary in a systematical way with ISM properties. A particularly strong trend is the decrease of the normalization of the relation with a lowering of the dust-to-gas ratio of the ISM. We show that this trend and other properties of the pressure -- $\\H2$ relation are natural consequences of the transition from atomic to molecular hydrogen with gas surface density. ",
        "introduction": "\\label{sect:intro}  The abundance of molecular gas in galaxies is set by the complex interplay of various formation and destruction processes operating in a highly turbulent medium. While many of the individual physical mechanisms are relatively well understood, such as the formation of molecular hydrogen on dust grains, its photo-dissociation by ultraviolet (UV) photons in the Lyman-Werner bands, or the importance of dust and $\\H2$ self-shielding, e.g., \\cite{1978ApJS...36..595D, 1979ApJS...41..555H, 1988ApJ...332..400S, 1996ApJ...468..269D}, we still lack a realistic and coherent picture that links together the formation of molecular gas, star formation, the turbulent, multi-phase structure of the ISM, and the importance of the various feedback channels due to star formation. The molecular content of galaxies is a key diagnostic that provides insights not only into how galaxies evolve and grow their stellar component, but also helps to constrain the properties of the physical processes (including feedback) that operate in the ISM. In addition, the modeling of the molecular ISM provides a crucial theoretical background for the interpretation of molecular gas surveys, such as those expected in the near future with the Atacama Large Millimeter/sub-millimeter Array (ALMA).  In addition to analytical and numerical models (\\citealt{2006ApJ...645.1024P, 2007ApJ...659.1317G, 2008ApJ...680.1083R, 2008ApJ...689..865K, 2009ApJ...693..216K, 2009ApJ...697...55G, 2010ApJ...721..975O, 2011ApJ...728...88G}), empirical correlations inferred from observations of nearby galaxies are often used to predict the molecular gas abundance. In particular, the correlation between the surface density ratio of molecular and atomic ($\\HI$) hydrogen $R_{\\rm mol}=\\Sigma_\\H2/\\Sigma_\\HI$ and the mid-plane pressure $P_{\\rm ext}$ (\\citealt{2002ApJ...569..157W, 2004ApJ...612L..29B, 2006ApJ...650..933B}) has been included in semi-analytical models (e.g., \\citealt{2009MNRAS.396..141D, 2010MNRAS.409..515F}) and numerical simulations (e.g., \\citealt{2010MNRAS.405.1491M}) to estimate $\\H2$ mass fractions of galaxies and their star formation rates, or in order to predict the global evolution of the $\\H2$ baryon content in the universe \\citep{2009ApJ...696L.129O}. Implicit in this approach is the assumption that the empirical correlation continues to hold for galaxies with ISM properties that are potentially different from those found in galaxies in the local universe. It is clearly crucial to test this assumption, either observationally \\citep{2010ApJ...722..919F}, or, as is the approach of this paper, with the help of numerical simulations that are based on a firm theoretical modeling of the microphysics of the ISM.  A further complication is the fact that $\\H2$ masses and surface densities are typically not directly accessible to observations. The kinetic temperature ($\\sim{}10$ K) of the bulk of the molecular hydrogen in galaxies is too low, and the gas sufficiently shielded from UV radiation, to populate the excited levels of the rotational ladder at a significant level \\citep{1982ARA&A..20..163S}. Hence, $\\H2$ masses are often inferred from the emission of  tracer elements and molecules. In particular, the optically thick emission from the main isotope of carbon-monoxide (CO) serves as a relatively reliable tracer of $\\H2$ mass, at least under conditions typical of molecular clouds in the Milky Way \\citep{1975ApJ...202...50D, 1986ApJ...309..326D, 1986A&A...154...25B, 1988ApJ...325..389M, 1994ApJ...429..694L, 1996A&A...308L..21S, 1997ApJ...481..205H, 2001ApJ...547..792D, 2007ApJ...663..866D, 2010ApJ...710..133A, 2012arXiv1202.4039T}. However, the conversion factor between $\\CO$ luminosity and $\\H2$ mass is expect to change systematically with the dust-to-gas ratio and with the strength of the interstellar radiation field \\citep{1995ApJ...448L..97W, 1996PASJ...48..275A, 2002A&A...384...33B, 2005A&A...438..855I, 2011MNRAS.412..337G, 2011ApJ...731...25K, 2011MNRAS.412.1686S, 2011ApJ...737...12L, 2012ApJ...746...69G, 2012ApJ...747..124F, 2012MNRAS.421.3127N}. This effect needs to be included in the numerical modeling of empirical relations that are based on CO observations.  In this paper we use numerical simulations of galaxies in a cosmological framework to study the origin of the  $P_{\\rm ext} - R_{\\rm mol}$ relation and its dependence on the properties of the ISM. Our numerical models compute the local $\\HI$ and $\\H2$ abundances in the ISM based on a chemical network of well understood formation and destruction processes. In addition, we compute the CO emission expected from our model galaxies to account for variations of the $\\CO-\\H2$ conversion factor. We show that, under Milky Way like ISM conditions, a $P_{\\rm ext} - R_{\\rm mol}$ relation similar to that seen in nearby galaxies arises rather naturally. We further show that galaxies with different dust-to-gas ratios and/or radiation fields also follow a $P_{\\rm ext} - R_{\\rm mol}$ relation, but with changes in the normalization and the slope. We conclude that $\\H2$ abundances estimated from a $P_{\\rm ext} - R_{\\rm mol}$ relation are not robust if these changes are not taken into account.  The outline of the paper is as follows. In section \\ref{sect:methods} we describe the set-up of our numerical approach, discuss details of the data analysis, and introduce the observational data sets that we use to compare with our numerical models. In the subsequent section \\ref{sect:results} we present our numerical predictions for the $P_{\\rm ext} - R_{\\rm mol}$ relation and its dependence on environmental parameters. We also discuss the role of the $\\CO$-$\\H2$ conversion factor and present a physical model that captures many of the properties of the $P_{\\rm ext} - R_{\\rm mol}$ relation. We summarize our results and conclude in section \\ref{sect:summary}. ",
        "conclusions": "\\label{sect:summary}  We studied the relation between mid-plane pressure and the $\\Sigma_\\H2/\\Sigma_\\HI$ ratio in a set of hydrodynamical simulations that follow explicitly the formation and destruction of molecular hydrogen in the ISM. We have demonstrated that these simulation predict a $P_{\\rm ext}-R_{\\rm mol}$ relation that is very similar to the one observed in nearby galaxies. The fact that the relation changes systematically with the dust-to-gas ratio and the strength of the UV radiation field in the ISM indicates that the $P_{\\rm ext}-R_{\\rm mol}$ relation should not be used as a universal tool to estimate $\\H2$ abundances in galaxies with ISM conditions different from those in the Milky Way. In particular, we find that the normalization of the $\\log_{10} P_{\\rm ext} - \\log_{10} R_{\\rm mol}$ relation decreases with a decreasing dust-to-gas ratio $\\DMW$. The scaling of the normalization is $\\propto\\DMW$ as long as the ISM regions that are included in fitting the $\\log_{10} P_{\\rm ext} - \\log_{10} R_{\\rm mol}$ relation are predominantly molecular and the metallicity dependence of the $\\CO$-$\\H2$ conversion factor is taken into account. We have proposed a simple model that is based on the numerical study of the $\\HI$ to $\\H2$ transition that captures many of the basic properties of the $P_{\\rm ext}-R_{\\rm mol}$ relation.  \\cite{2008ApJ...680.1083R} studied the $P_{\\rm ext}-R_{\\rm mol}$ relation in hydrodynamical simulations of isolated disk galaxies. While their numerical approach differs in detail, they also find a clear correlation between $P_{\\rm ext}$ and $R_{\\rm mol}$ with a slope $\\sim{}0.9$. Although most ISM regions included in their analysis probe the $\\HI$ dominated regime, their largest galaxy reaches pressures high enough to study the $P_{\\rm ext}-R_{\\rm mol}$ relation in molecular hydrogen dominated gas. Interestingly, while our numerical modeling predicts a change in the slope of the relation from $\\sim{}1.3$ for $R_{\\rm mol}\\ll{}1$ to $\\sim{}0.7$ for $R_{\\rm mol}\\gg{}1$, such a change is not immediately apparent in the work by \\cite{2008ApJ...680.1083R}. As showed in the preceding section, a slope of $\\lesssim{}0.7$ in the $\\H2$ dominated regime is a direct consequence of a fixed $\\HI$ saturation limit. Hence, the relatively steep slope ($>0.7$) found by \\cite{2008ApJ...680.1083R} in the $\\H2$ dominated regime might indicate that in their numerical model $\\Sigma_\\HI$ decreases with increasing $\\Sigma_g$ at large gas surface densities.  \\cite{1993ApJ...411..170E} studied the origin of the relation between $\\H2$ abundance and mid-plane pressure analytically. He found that in the $\\HI$ dominated regime $R_{\\rm mol}\\propto{}f_\\H2\\propto{}P_{\\rm ext}^{2.2}/j$, where $j$ is the local radiation field. While this relation is sometimes used to ``explain'' the $P_{\\rm ext}-R_{\\rm mol}$ relation, we stress that the $\\HI$ dominated regime is not the regime typically studied in observations. Furthermore, $j$, the interstellar radiation field incident on the clouds in the ISM, is likely highly spatially variable and will correlate only to a certain degree with galactic properties smoothed over $\\sim{}$kpc scales. This implies that applying this formula properly is non-trivial. For instance, if the $\\sim{}$kpc scale UV radiation field is approximately constant a naive application of Elmegreen's formula predicts $R_{\\rm mol}\\propto{}P_{\\rm ext}^{2.2}$. In contrast, the model presented in section \\ref{sect:model} predicts $R_{\\rm mol}\\propto{}P_{\\rm ext}^{4/(2+\\beta{})}$ for $R_{\\rm mol}<1$, which results in a scaling $R_{\\rm mol}\\propto{}P_{\\rm ext}^{1-1.33}$ (for $\\beta{}\\sim{}1-2$) close to observations \\citep{2008AJ....136.2782L}.  The numerical predictions presented in this paper are based on the modeling of a variety of well understood physical processes in the ISM. However, it is possible or even likely that we are missing processes that have an impact on the atomic and molecular hydrogen abundance in galaxies. For instance, while the simulations include certain aspects of stellar feedback (e.g., ionizing and dissociating radiation from massive stars, metal enrichment from supernovae, stellar mass loss), they miss others (e.g., radiation pressure, thermal feedback from supernovae, stellar winds). Also, the resolution of the simulations is too coarse to resolve individual molecular clouds and the detailed dynamics and small scale chemistry that takes place within them. Hence, it will be important to verify the predictions of our numerical models observationally. Important tests include (i) the measurement of the normalization of the $\\log_{10} P_{\\rm ext} - \\log_{10} R_{\\rm mol}$ relation in galaxies with low metallicities and presumably low dust-to-gas ratios, (ii) the measurement of the slope of the relation in both the $\\HI$ and $\\H2$ dominated region, and (iii) to check for correlations of the normalization and slope with the scaling exponent between the stellar surface density and the gas surface density."
    },
    "1208/1208.3206_arXiv.txt": {
        "abstract": "{In the last decades cosmological N-body dark matter simulations have enabled ab initio studies of the formation of structure in the Universe. Gravity amplified small density fluctuations generated shortly after the Big Bang, leading to the formation of galaxies in the cosmic web. These calculations have led to a growing demand for methods to analyze time-dependent particle based simulations. Rendering methods for such N-body simulation data usually employ some kind of splatting approach via point based rendering primitives and approximate the spatial distributions of physical quantities using kernel interpolation techniques,  common in SPH ({\\sl Smoothed Particle Hydrodynamics})-codes. This paper proposes three GPU-assisted rendering approaches, based on a new, more accurate method to compute the physical densities of dark matter simulation data. It uses full phase-space information to generate a tetrahedral tessellation of the computational domain, with mesh vertices defined by the simulation's dark matter particle positions. Over time the mesh is deformed by gravitational forces, causing the tetrahedral cells to warp and overlap. The new methods are well suited to visualize the cosmic web. In particular they preserve caustics, regions of high density that emerge, when several streams of dark matter particles share the same location in space, indicating the formation of structures like sheets, filaments and halos. We demonstrate the superior image quality of the new approaches in a comparison with three standard rendering techniques for N-body simulation data. } ",
        "introduction": " ",
        "conclusions": "\\label{Sec-Conclusions}  We presented three GPU-accelerated rendering approaches for N-body dark matter simulation data, based on a tetrahedral decomposition of the computational domain that allows a physically more accurate estimation of the mass density between the tracer particles than previous methods. They use the full phase-space information of the ensemble of dark matter tracer particles and two of them minimize pre-processing time (centroids and cell-projection) as well as data transfer between the CPU and GPU, by generating all connectivity information as well as the derived quantities, like mass density of the tetrahedral mesh elements, on the GPU. Thus these approaches are particularly well suited for time-dependent data. Their performance should benefit significantly from the increased number of cores expected for future generations of graphics hardware.  We compared these new methods to three standard rendering approaches for dark matter simulations: two  based on constant and adaptive kernel sizes that estimate the local densities from the nearest-neighbors, as well as a Voronoi tessellation generated by the simulations tracer particles. We showed that our approaches yield considerably better image quality with less pre-processing times and graphics memory requirements. The full tetrahedral cell-projection methods clearly stands apart, however. Without artificial smoothing or density estimates derived from averaging over the particle distribution, features previously washed out, become clearly visible and give new insight in the physical large-scale features of the {\\sl cosmic web}, including voids, filaments and halos."
    },
    "1208/1208.3947_arXiv.txt": {
        "abstract": "One of the most striking aspects of the 11-year sunspot cycle is that there have been times in the past when some cycles went missing, a most well-known example of this being the Maunder minimum during 1645---1715. Analyses of cosmogenic isotopes ($^{14}$C and $^{10}$Be) indicated that there were about 27 grand minima in the last 11,000~yr, implying that about 2.7\\% of the solar cycles had conditions appropriate for forcing the Sun into grand minima. We address the question how grand minima are produced and specifically calculate the frequency of occurrence of grand minima from a theoretical dynamo model. We assume that fluctuations in the poloidal field generation mechanism and in the meridional circulation produce irregularities of sunspot cycles. Taking these fluctuations to be Gaussian and estimating the values of important parameters from the data of last 28 solar cycles, we show from our flux transport dynamo model that about 1--4\\% of the sunspot cycles may have conditions suitable for inducing grand minima. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.1490_arXiv.txt": {
        "abstract": "In the present work we expand the study of time-dependent ionization previously identified to be of pivotal importance for acoustic waves in solar magnetic flux tube simulations.  We focus on longitudinal tube waves (LTW) known to be an important heating agent of solar magnetic regions. Our models also consider new results of wave energy generation as well as an updated determination of the mixing length of convection now identified as 1.8 scale heights in the upper solar convective layers.  We present 1-D wave simulations for the solar chromosphere by studying tubes of different spreading as function of height aimed at representing tubes in environments of different magnetic filling factors.  Multi-level radiative transfer has been applied to correctly represent the total chromospheric emission function.  The effects of time-dependent ionization are significant in all models studied.  They are most pronounced behind strong shocks and in low density regions, i.e., the middle and high chromosphere. Concerning our models of different tube spreading, we attained pronounced differences between the various types of models, which were largely initiated by different degrees of dilution of the wave energy flux as well as the density structure partially shaped by strong shocks, if existing.  Models showing a quasi-steady rise of temperature with height are obtained via monochromatic waves akin to previous acoustic simulations.  However, longitudinal flux tube waves are identified as insufficient to heat the solar transition region and corona in agreement with previous studies. ",
        "introduction": "The outer atmospheric activity of the Sun as well as other late-type stars is largely determined by the structure and time evolution of photospheric magnetic fields.  These fields extend into the stellar outer atmosphere, where they cause nonradiative energy to be deposited \\citep[e.g.,][]{and89,lin91,schr01,bog03,rob04,ulm04,mus04}. This energy is distributed over the chromosphere, transition region, and corona; it is also considered pivotal for the heating and acceleration of solar and stellar winds.  For new observations and simulations of the relative importance of acoustic and magnetic energy deposition in the solar chromosphere, see, e.g., \\cite{fos05}, \\citet*{cun07}, \\cite{kal07}, and \\cite{bel09}.  Recent joint observational and theoretical studies aimed at elucidating the significance of wave heating in the solar photosphere and lower chromosphere were given by \\cite{bec10}, \\citet*{bec12}, and \\cite{wed12}.  The latter study investigates the channeling of magnetic energy through the solar chromosphere into the corona, thus resembling super-tornadoes under solar conditions. Notable reviews aimed at solar chromospheric and coronal heating were given by \\cite{nar96}, \\cite{kli06}, and \\cite{erd07}.  Magnetic flux tubes are an important feature of the solar surface structure \\citep[e.g.,][]{spr83,sol96,schr00}.  Both observational and theoretical studies showed that they are carriers of longitudinal tube waves \\cite[e.g.,][]{her85,sol92,has08}, which give rise to considerable temperature increases as function of height as revealed by chromospheric spectral features. In this paper, we continue to pursue our previous line of research focused on the generation of different types of waves \\citep[e.g.,][]{nar96}, effects of the propagation and dissipation of waves \\citep*[e.g.,][]{her85,faw98,cun99}, and the emergence of chromospheric emission \\citep[e.g.,][]{cun99,faw02,ramc03}. However, in accord with previous simulations of acoustic waves \\citep[e.g.,][]{car92,car95,ramu03,ramc05b,cun07}, the described longitudinal wave simulations will also employ time-dependent (i.e., non-instantaneous) ionization.  The geometrical and thermodynamic properties of the tube atmospheres are expected to also impact the dissipation of the wave energy as well as the so-called energy velocity as recently pointed out by \\cite{wor12}.  Time-dependent ionization entails that, e.g., behind strong shocks the long timescales of hydrogen ionization / recombination initially prevent the dissipated energy to be converted into ionization energy, thus leading to strong temerature spikes as well as a variety of other dynamic phenomena.  There exists a great motivation to revisit the dissipation of acoustic and magnetic waves in the solar chromosphere, which is the new determination of the mixing length near the top of the solar convective zones.  \\cite{stei09a,stei09b} who provided new state-of-the-art simulations of the solar convection zone on the scale of supergranules, extending 10\\% of its depth but half of its pressure scale height, deduced a mass mixing length of $\\alpha_{\\rm ML} = 1.8$, thus superseding the previous results of 1.5 or 2.0 \\citep{stef93,tra97}, which were widely used in previous solar heating computations.  In our study we will also consider the relevance of tube spreading, i.e., different tube opening radii, for the energetics and thermodynamics of the magnetically heated solar chromospheric structure.  Early results based on adiabatic longitudinal waves without the consideration of time-dependent ionization have been given by \\cite{faw98}.  They found that the tube shape is of critical importance for the heating of flux tubes.  In fact, tubes of wide opening radii show little heating, whereas constant cross section tubes show very large heating at all heights. Previous simulations of longitudinal waves for solar coronal tube structure have been given by, e.g., \\citet*{ofm99}, \\citet*{ofm00}, and \\cite{cuns04}. These results further highlight the importance of tube spreading with respect to the shock wave amplitude and time-dependent and height-dependent heating rate.  Concerning the formation of Ca~II and Mg~II emission, we will consider a two-component (magnetic and acoustic) model of the solar chromosphere with heating by longitudinal tube waves inside the flux tubes and heating by acoustic waves outside the flux tubes.  The governing equations as well as the methods of our study, including the computation of the initial flux tube models, are discussed in Sect.~2.  In Sect.~3, we describe the results of our longtitudinal wave simulations for solar flux tube models with different tube spreadings; the latter correspond to different magnetic filling factors.  The focus of these studies concerns the effects of time-dependent ionization.  Our summary and conclusions are given in Sect.~4. ",
        "conclusions": "We pursued sets of time-dependent model simulations for solar magnetic flux tubes with focus on longitudinal tube waves with and without the consideration of time-dependent ionization, particularly pertaining to hydrogen.  The treatment of longitudinal tube waves is motivated by a large variety of studies including work by \\cite{has08} arguing that longitudinal tube waves are able to provide quasi-steady heating sufficient to explain the bright solar network grains observed in Ca~II H and K.  We studied the dynamics of flux tubes with different magnetic filling factors, i.e., 0.1\\%, 0.01\\%, and 0.001\\%, corresponding to different tube top opening radii, guided and motivated by previous observational results \\citep[e.g.,][]{spr83,sol96}.  It was found that the dynamic tube structures are determined by a complex interplay between shock heating, radiative and hydrodynamic cooling (with the latter caused by wave pressure), dilution of the wave energy flux due to the flux tube geometry, and time-dependent ionization. We identified pronounced differences between the tubes models especially in regard to tubes with different tube opening radii and in response to the inclusion or omission of time-dependent ionization.  Similar to the case of acoustic waves, time-dependent ionization in LTW models leads to a large range of phenomena, such as over- and underionization of the flow, increased temperature jumps at shocks, and modified mean (i.e., time-averaged) temperature, density and ionization structures.  The impact of time-dependent treatment of ionization for flux tubes becomes evident through comparisons of mean temperatures between the different types of models.  Regarding monochromatic models it was found that time-dependent ionization leads to lower average temperatures inside flux tubes compared to time-independent ionization, an effect already known for acoustically heated models.  However, the average temperatures inside of flux tubes in models with and without time-dependent ionization are nonetheless noticeably higher than in the acoustically heated external atmosphere. If monochromatic longitudinal tube waves are used, the mean (i.e., time-averaged) temperatures of the tubes generally show a quasi-steady increase with height.  However, if frequency spectra are used, the mean temperatures, in essence, show no rise with height at all. This behaviour is initiated by the formation of very strong shocks (especially in narrow tubes, corresponding to small magnetic filling factors) resulting in large-scale quasi-adiabatic cooling, leading to unrealistically low temperatures. This type of result has already previously been found in corresponding 1-D models of acoustic waves \\cite[e.g.,][]{car92,car95,ramu03}, a behaviour deemed highly unrealistic \\citep{ulm05}.  An appropriate enhancement of our current models will be the calculation of self-consistent 3-D magnetohydrodynmic (MHD) models. Those models face, however, the principal challenge of the necessity to include both detailed multi-level 3-D radiative transfer and 3-D flows including the detailed treatment of shock formation and shock interaction.  Important progress has already been made \\citep[e.g.,][]{stei09a,stei09b}, but further efforts are needed to obtain of a concise picture.  Tentative insights into the principal properties of these types of future models can be attained through inspecting existing 3-D time-dependent {\\it non-magnetic} hydrodynamics models, which also consider time-dependent non-equilibrium effects caused by hydrogen \\citep{lee06}, albeit various restrictive assumptions including (but not limited to) the lack of back-coupling of the ionization to the equation of state.  In this type of models it is found that the build-up of strong shocks due to shock interaction is largely absent, resulting in a lack of unrealistically high cooling behind the strong shocks previously also referred to as ``hydrodynamic refrigeration\" \\citep{cunm94}.  In this case a quasi-steady rise of temperature with height is attained, which appears to be in close resemblance to empirical solar chromosphere models \\citep[e.g.,][]{and89}.  Nonetheless, our results based on time-dependent ionization and LTW frequency spectra allow insight into the limiting case of 1-D geometry, while also noting that our models based on time-dependent ionization and monochromatic waves are expected to be approximately reflective of physical reality.  An alternative, or perhaps supplementary, way of supplying chromospheric heating might be given through ambipolar diffusion as described by \\cite{kho12}.  Here the presence of neutrals, together with the decrease with height of the collisional coupling, leads to deviations from the classical magnetohydrodynamic behaviour of the chromospheric plasma.  \\cite{kho12} pointed out that a relative net motion occurs between the neutral and ionized components, referred to as ambipolar diffusion. According to this model, the dissipation of currents in the chromosphere is enhanced by orders of magnitude due to the action of ambipolar diffusion, as compared with the standard ohmic diffusion.  The authors proposed that a significant amount of magnetic energy can be released to the chromosphere just by existing force-free 10--40~G magnetic fields there.  Additional studies were given by \\cite*{fed09}, \\cite*{vig09}, and \\cite{erd10}.  \\cite{fed09} studied the oscillatory response of the 3-D solar photosphere to the leakage of photospheric motion.  They found, among other results, that high-frequency waves propagate from the lower atmosphere across the transition region experiencing relatively low reflection, and transmitting most of their energy into the corona, and, furthermore, that the magnetic field acts as a waveguide for both high- and low-frequency waves originating from the photosphere and propagating up into the solar corona.  \\cite{vig09} provided a targeted study on wave propagation and energy transport in the magnetic network of the Sun based on 2-D MHD simulations, which among other results identified the limited capacity of acoustic waves.  \\cite{erd10} investigated the oscillatory modes of a magnetically twisted compressible flux tube embedded in a compressible magnetic setting, including applications to solar magneto-seismology.  As an overarching statement concerning our study we conclude that the significance of time-dependent ionization identified in the simulations of longitudinal tube waves is a stark motivation to also consider this type of effect in future models of transverse and torsional tube waves.  This will allow to obtain a more detailed picture of the dynamics and energetics of solar-type outer atmospheres.  This obvious suggestion is also supported by the repeatedly obtained finding that longitudinal flux tube waves, as gauged through models considering time-dependent ionization phenomena as done in the present study, are insufficent to supply an adequte amount of energy for balancing coronal heating in the view of early and updated estimates by \\cite{gue07} and others."
    },
    "1208/1208.2493_arXiv.txt": {
        "abstract": "SVOM (Space-based multi-band astronomical Variable Object Monitor) is a future Chinese-French satellite mission which is dedicated to Gamma-Ray Burst (GRB) studies. Its anti-solar pointing strategy makes the Earth cross the field of view of its payload every orbit. In this paper, we present the variations of the gamma-ray background of the two high energy instruments aboard SVOM, the Gamma-Ray Monitor (GRM) and ECLAIRs, as a function of the Earth position. We conclude with an estimate of the Earth influence on their sensitivity and their GRB detection capability. ",
        "introduction": "\\label{sec:intro} After forty years of studies, we are still far from a full understanding of Gamma-Ray Bursts (GRBs). Accurate measurements of GRB parameters, such as their position, redshift, peak-energy and so on, are needed to further understand GRBs themselves and their use as astrophysical tools. SVOM (Space-based multi-band astronomical Variable Object Monitor) \\cite{Paul2011}\\cite{Basa2008} is a future Chinese-French satellite mission which is dedicated to GRB studies. The payload aboard SVOM consists of two high-energy instruments, the GRM and ECLAIRs, together with two low-energy instruments. GRM is a Gamma Ray Monitor, and ECLAIRs is a coded-mask camera for X- and Gamma-rays. Together, they can provide GRB localizations and GRB spectral observations from a few keV to a few MeV. This wide spectral range is crucial to determine as accurately as possible the GRB peak-energy, which is a key parameter to make use of GRBs as \u201cstandard candles\u201d and extend the measurement of the cosmological parameters at large redshifts \\cite{Paul2011}\\cite{Ghirlanda2006}.  SVOM is a Low Earth Orbit (LEO) mission with an altitude of 600 km and an inclination of 30$^\\circ$. Its instruments point close the anti-solar direction during a large fraction of the orbit in order to permit the follow-up observations by large ground-based telescopes to optimize the redshift measurements \\cite{Paul2011}\\cite{Cordier2008}. This strategy makes the Earth get in and out of the detector\u2019s Field of the View (thereafter FoV) once per orbit. The detectors will not operate when crossing the South Atlantic Anomaly (SAA) for protection from the too high particle flux, which results in a decrease of the observation time. These factors will cause the instrumental background to change when the Earth gets in and out of the FoV and with time. In this paper, we present a study on how the gamma-ray background of GRM and ECLAIRs changes with the Earth position in the FoV. The impact of these changes on the sensitivity and GRB detection rate are estimated as well. This background research laid the groundwork for the investigation of the GRB trigger strategy of GRM and ECLAIRs. This paper is organized as follows:  \\S\\ref{sec:instr} In this section, we briefly describe the two high energy instruments aboard SVOM.  \\S\\ref{sec:simu} The detector mass-models, the simulation method and the spectral models of gamma-ray background sources are described in this section.  \\S\\ref{sec:results} The background simulation results, the computations of the detector sensitivities and the GRB detection rates are presented in this section.  \\S\\ref{sec:conc} We conclude with a concise summary and a discussion of some limitations of our method as well as some further improvements. ",
        "conclusions": "\\label{sec:conc} In this work, three kinds of background varying with Earth positions have been simulated for GRM and ECLAIRs which are two high-energy instruments aboard the SVOM mission. In particularly, the simulations of ALBEDO with the model of Sazonov2007 are described in detail for its complexity and precision. Following the background simulations, the corresponding sensitivities and GRB detection rates are calculated based on the given observations.  Concerning the impact of each background component on the instruments, generally, for both GRM\\_NaI and ECLAIRs, the cosmic gamma-rays are the most important background source when the Earth is outside of the FoV, especially in the relatively low energy range. The albedo gamma-rays are always the most important background source in the relatively high energy range. The reflected cosmic gamma-rays, which are lower than cosmic gamma-rays in flux and lower than albedo gamma-rays in energy, have the least impact on the instruments, but are still an important component in the low energy range when the Earth enters the FoV. For GRM\\_CsI, albedo gamma-rays are always the dominant background source.  As for the spectral models used in the simulations, the limitation lays in the energy range. ECLAIRs is sensitive from 4 keV to 250 keV and GRM is sensitive from 30 keV to 5 MeV. However, the model of Moretti2009 was obtained considering the data in the energy range 1.5-200 keV. The model of Sazonov2007 is a function of the energy which covers from 25 keV to 300 keV. Only the energy range of the Churazov2006 model reaches the MeV range. As a result, one further work on the gamma-ray background estimations is to find better source models for the wide energy range of SVOM. Maybe different models should be used for different energy ranges.  In the calculations in Section \\ref{sec:sensitivity}, only gamma-ray background components are considered without including other kinds of backgrounds induced by protons, electrons, positrons, etc. as well as the background produced by the instrument itself. However, we cannot conclude that the value of the GRB detection rate is overestimated for this reason. Because there are many factors affecting the final results in fact. For example, to compute the sensitivity values in Table 2 accurately, the efficiency corresponding to different incident directions of photons from the FoV should be considered instead of only the vertical direction. Furthermore, the sensitivity used to trigger GRBs on board will consider diverse energy ranges, integral time scales and significance levels. One assumption in our computation is that the satellite has seven states relative to the Earth positions and each state lasts 1/7 of the satellite orbit period. This is a simplification, since the real pointing law on orbit will be more complex. Back to our gamma-ray background simulations, the maximum ALBEDO spectrum with the minimum solar modulation 0.25 GV at the highest level orbit position (30$^\\circ$N, 287$^\\circ$E) was used. Accordingly, a more accurate estimation of the GRB detection rate needs further work and another article will be prepared."
    },
    "1208/1208.2988_arXiv.txt": {
        "abstract": "{The $Kepler$ object KIC 12557548 b is peculiar. It exhibits transit-like features every 15.7 hours that vary in depth between 0.2\\% and 1.2\\%. Rappaport et al. (2012) explain the observations in terms of a disintegrating, rocky planet that has a trailing cloud of dust created and constantly replenished by thermal surface erosion. The variability of the transit depth is then a consequence of changes in the cloud optical depth.} {We aim to validate the disintegrating-planet scenario by modeling the detailed shape of the observed light curve, and thereby constrain the cloud particle properties to better understand the nature of this intriguing object.} {We analysed the six publicly-available quarters of raw $Kepler$ data, phase-folded the light curve and fitted it to a model for the trailing dust cloud. Constraints on the particle properties were investigated with a light-scattering code.} {The light curve exhibits clear signatures of light scattering and absorption by dust, including a brightening in flux just before ingress correlated with the transit depth and explained by forward scattering, and an asymmetry in the transit light curve shape, which is easily reproduced by an exponentially decaying distribution of optically thin dust, with a typical grain size of 0.1~$\\mu$m.} {\\rm Our quantitative analysis supports the hypothesis that the transit signal of KIC 12557548 b is due to a variable cloud of dust, most likely originating from a disintegrating object.} ",
        "introduction": "\\citet{rap12} recently reported the discovery of a peculiar signal in the light curve of KIC 12557548\\footnote{Henceforth denoted with KIC1255 b}. A transit-like signal is seen every 15.7 hours, with a depth greatly varying from event to event, ranging from less than 0.2\\% to more than 1.2\\%. The transit timing is consistent with a single period and shows no variations down to the 10$^{-5}$ level. This suggests that the changes in transit depth cannot be produced by a fast precession of the orbit, which would in turn require the gravitational perturbation of a second companion in the system. In addition, the phase-folded light curve shows no evidence for ellipsoidal light variations down to 5 parts in 10$^5$. This sets an upper limit of about 3~$M_{Jup}$ on the mass of the orbiting object. The explanation proposed by \\citet{rap12} entails the disintegration of a super-Mercury, caused by the high surface temperatures. Material stripped away from the planet surface forms and stochastically replenishes a cloud of dust that obscures part of the stellar disk, causing the dimming of the stellar light to randomly vary from orbit to orbit. Only a planet with mass on the order of Mercury would allow dust particles to accelerate to the escape velocity and form a trailing cloud of sufficient size.  To test the interpretation of \\citet{rap12}, a quantitative understanding of the observed transit light curve is required, in particular the asymmetry of the in-transit portion of the light curve and the excess of flux before ingress. We reduced the complete set of publicly available {\\it Kepler} observations, phase-folded the light curve and grouped the transits based on their depth, as explained in Sect.~\\ref{data}. A model for the trailing dust, described in Sect.~\\ref{model}, is then fitted to the data to estimate the transit parameters, the basic shape of the cloud and the average size of the dust particles. The results of our analysis and their interpretation are discussed in Sect.~\\ref{results}. ",
        "conclusions": "A simple, one-dimensional model of a trailing dust cloud can reproduce the observed light curve of KIC1255 b in exquisite detail. Features such as the flux excess before ingress, the sharp drop in ingress, the pointed shape of the flux minimum and the long flux recovery after egress are all signs of a trailing cloud of dust that periodically occults the star. The most likely origin of this trailing dust is a planetary or sub-planetary body that evaporates due to the intense stellar irradiation. If this is indeed the case, we are presented with the unique opportunity to probe the interior of an exoplanet. Although the geometry of the transit is well constrained by the currently available {\\it Kepler} data, further observations with a faster cadence and multi-wavelength coverage are required to not only constrain the size but also the composition of the grains."
    },
    "1208/1208.1791.txt": {
        "abstract": "We present an analysis of the ionic composition of iron for two interplanetary coronal mass ejections observed in May 21-23 2007 by the ACE and STEREO spacecraft in the context of the magnetic structure of the ejecta flux rope, sheath region, and surrounding solar wind flow. This analysis is made possible due to recent advances in multispacecraft data interpolation, reconstruction, and visualization as well as results from recent modeling of ionic charge states in MHD simulations of magnetic breakout and flux cancellation CME initiation.  We use these advances to interpret specific features of the ICME plasma composition resulting from the magnetic topology and evolution of the CME. We find that in both the data and our MHD simulations, the flux ropes centers are relatively cool, while charge state enhancements surround and trail the flux ropes. % The magnetic orientation of the ICMEs are suggestive of magnetic breakout-like reconnection during the eruption process, which could explain the spatial location of the observed iron enhancements just outside the traditional flux rope magnetic signatures and between the two ICMEs. % Detailed comparisons between the simulations and data were more complicated, but a sharp increase in high iron charge states in the ACE and STEREO-A data during the second flux rope corresponds well to similar features in the flux cancellation results.  We discuss the prospects of this integrated in-situ data analysis and modeling approach to advancing our understanding of the unified CME-to-ICME evolution. ",
        "introduction": "Coronal mass ejections (CMEs) are periods of explosive magnetic energy release and coronal field reconfiguration that blow out huge portions of the quasi-stable solar atmosphere into interplanetary space \\citep{Tousey1973,Gosling1997}. When these solar eruptions are Earth-directed, they generally arrive in 1--5 days \\citep{Cane2003} and their impact can cause geoeffective space weather responses \\citep{Gosling1993, Zhang2004}. CMEs are observed soon after their eruption by remote sensing instruments that capture the global structure of these events. Their morphological structure in coronagraph observations typically consists of a bright bubble of plasma, sometimes with a darker cavity and bright core \\citep[e.g.,][]{Illing1986,Howard1997}. Once these CMEs propagate into the heliosphere they are commonly referred to as interplanetary CMEs (ICMEs) and are observed in-situ by spacecraft that sample local measurements of the magnetic field, plasma, and composition. ICME identification is based on a variety of parameters including enhancements in magnetic field,  composition and temperature depressions \\citep[e.g., see][and references therein]{Zurbuchen2006}.  Though CMEs and ICMEs are the same phenomena, the remote sensing and in-situ observations can often be difficult to directly relate to one another, due to both the nature of the observations as well as the evolution that takes place during the transit from the Sun to the spacecraft. In-situ solar wind composition data provides an important connection between the ICME observations and associated observations at the Sun, in particular providing a measure of physical properties in the corona that can be compared to EUV and X-ray emission and white light coronagraph observations. Charge state composition is directly related to the electron temperature at the source region, with hotter source material resulting in higher in-situ charge states \\citep{Buergi1986}. During the solar wind expansion, these charge states \u00d2freeze-in\u00d3 once the expansion timescale becomes larger than the collision timescale \\citep{Hundhausen1968}. {ICMEs often contain unusual charge state distributions. These charge states can be enhanced relative to the ambient solar wind due to heating at the source region, lower than the ambient solar wind due to incompletely heated cold prominence material, or a combination of both low and high charge states \\citep{Bame1979, Gloeckler1999, Richardson2004,  Lepri2010, Gruesbeck2011, Gilbert2012}. Because of this complexity, it is important to examine the full charge state distribution. Charge state enhancements are weakly correlated with associated flare magnitude \\citep{Reinard2005} and tend to be strongest near the center of the ICME \\citep{Reinard2008}. In this paper we examine the in-situ iron charge state data  from the ACE and STEREO spacecraft and compare them with theoretical predictions derived from MHD simulations of two different CME initiation models using the technique developed by \\citet{Lynch2011}.  Given the localized nature of ICME observations (ICMEs are typically observed by a single spacecraft that samples a narrow trajectory through the larger CME structure) drawing conclusions about the global structure of ICMEs can be complicated as it is not always clear which part of the ICME is being sampled. The subset of ICMEs with the most structured field and plasma signatures are known as magnetic clouds and are defined as having an enhanced field strength, relatively smooth magnetic field rotations, and a lower than expected temperature \\citep{Burlaga1981}. While the commonly accepted occurrence rate for magnetic clouds is between 30--50\\% of all ICME events \\citep{Gosling1990,Richardson2004}, there is increasing evidence that all or most CMEs develop or contain some sort of flux rope structure \\citep{Jian2006,Kilpua2011}. For those ICMEs, or portions of ICMEs, that meet the criteria to be considered magnetic clouds \\citep{Burlaga1981} the overall ICME structure can be approximated by modeling the in-situ ICME magnetic fields as force-free, cylindrical flux ropes \\citep[e.g.,][]{Lepping1990}. These techniques allow researchers to interpret the cavity portion of the CME coronagraph morphology with the in-situ magnetic flux rope structure \\citep[][]{Illing1986,Wood1999,Webb2009,Vourlidas2012} and indeed observations have confirmed this association \\citep[e.g.][]{DeForest2011}. However, analysis of the current density structure, measures of the interplanetary pressure-balance, and multispacecraft observations have shown that even magnetic clouds are not precisely force-free, cylindrical objects. Therefore non-force free and consequently non-cylindrical fits were developed \\citep[e.g.,][]{Farrugia1993,Farrugia1999,Mulligan2001,Hidalgo2002} that allowed more complexity in the local ICME field modeling, while still retaining a big picture view of the global structure. Though these advances have allowed a greater number of ICMEs to be modeled, all of these analytic flux rope modeling techniques are limited in that they are only relevant for the flux rope-like portion of the ICME; non-flux rope ICMEs or cuts through an ICME that do not cleanly intersect the flux rope cannot be included because they lack the proper magnetic structure needed for the fits. In theory, the pressure-balance Grad-Shafranov reconstruction techniques \\citep[e.g.,][]{Hu2001,Hu2002,Isavnin2011} avoid the a-priori flux rope criteria, as well as the problematic identification of the flux rope boundaries, but in practice are most successful when reconstructing flux rope structures \\citep{Riley2004,Al-Haddad2011}.  Recently, \\citet{Mulligan2012} addressed this issue by developing a technique to incorporate non-MC (magnetic cloud) material into MC flux rope modeling.  Multispacecraft observations of ICMEs have made and continue to make a valuable contribution to our understanding of the larger heliospheric structure, placing important constraints on flux rope modeling and the spatial extent of ICME ejecta over which the coherent flux rope signatures persist \\citep{Burlaga1981,Hammond1995,Mulligan1999,Mulligan2001,Riley2003}. The Solar-Terrestrial Observatory (STEREO) mission \\citep{Kaiser2008} was launched in October 2006 and consists of twin spacecraft at 1 AU traveling ahead of and behind the Earth at a rate of approximately 22.5 degrees per year. The STEREO observing geometry provides a means of relating multispacecraft in-situ observations to each other and to the continuous remote coronal and heliospheric imaging starting at the Sun and extending to 1 AU. The data from the SECCHI suite \\citep{Howard2008} combined with the IMPACT \\citep{Luhmann2008} and PLASTIC \\citep{Galvin2008} in-situ field and plasma measurements have ushered in a new era of direct CME-ICME observations, explicitly linking various morphological features to their associated in-situ properties \\citep[e.g.,][]{Harrison2008,Harrison2010,Davis2009,Davies2009,Moestl2009a,Roulliard2009,Webb2009,Lynch2010,DeForest2011}. In addition, the two STEREO spacecraft, along with the ACE and WIND spacecraft, offer three or four tracks through a given ICME rather than just one, allowing more of the ICME spatial structure to be sampled \\citep{Kilpua2011}. However, often even three tracks of observations through an ICME cannot fully constrain the ICME global structure and further heliospheric modeling is necessary to interpret and relate the multispacecraft data back to remote observations at the Sun. \\citet{Mulligan2012} describe a sophisticated method that provides an interpolated spatial mapping of the data between the spacecraft tracks. This method provides an estimate of the global ICME structure and simplifies the comparison with modeled and observed CME structures near the Sun.  The structure of the paper is as follows. In Section 2 we describe our methodology. First, in section 2.1, we briefly review previous work on the source region and the multispacecraft in-situ data of the 2007 May 21--23 ICMEs. In section 2.2 we describe the \\citet{Mulligan2012} reconstruction and interpolation techniques for transforming the in-situ timeseries data into spatial maps and present two-dimensional composition data in the context of the ICME flux rope structures and their surrounding regions. In section 2.3 we describe the \\citet{Lynch2011} application of ionic charge state calculations to numerical MHD simulations to derive theoretical spatial distributions of iron charge states for two idealized CME initiation scenarios. In section 3 we compare the observed and simulated ionic composition structure; first, on the global scale of the two flux rope ICMEs and their  spatial structure, and second, by constructing a set of synthetic spacecraft trajectories through the MHD simulation data for direct comparison with the multispacecraft measurements of the detailed iron charge state distributions (Fe$^{+6}$ to Fe$^{+20}$). In section 4 we discuss how the structure and complexity of the interplanetary composition data is reflected in the complexity of the May 19 and 20 CME source region's magnetic topology and resulting eruption scenario.  We conclude, in section 5, with a summary of our results and prospects for future analyses. ",
        "conclusions": "We have described an effort to understand the full global structure of ICMEs from the Sun to the Earth by combining a novel method to derive ionic charge state distributions from MHD models \\citep{Lynch2011} with newly developed, groundbreaking multispacecraft data analysis techniques that allow us to construct the global structure of ICMEs \\citep{Mulligan2012}. We compare the iron charge states derived from idealized MHD simulations with observations of the May 21-23, 2007 ICMEs. %   The large scale comparisons (Figures 3 and 4) show similarities between the $Q_{\\rm Fe}$ derived from MHD simulation results and the spatial mapping results derived from STEREO and ACE observations of the May 2007 events. In particular, we find that the charge state enhancements tend to occur at the back end and surrounding the flux rope material. Both the BM and the FC models have similar behavior with enhanced charge states surrounding and following the colder flux rope center. % Enhancements seen beyond the flux rope in the data, particularly ahead of the MC2 ICME, are well matched to the BM results, suggesting MC1 eruption's eruptive flare reconnection was simultaneously acting as breakout-like reconnection during the sympathetic eruption scenario resulting from the source region's multipolar topology.  More detailed time series comparisons between the spacecraft observations and cuts through the MHD model results (Figures 6 and 7) show that the ambient solar wind charge state levels are lower than predicted by the BM and higher than predicted by the FC model simulations. % Quantitative agreement between the observed flux rope iron charge state distributions are not well captured in the modeling results, particularly for MC1. However, the FC model does remarkably well at qualitatively matching the increased charge states observed in MC2, including both the $Q_{\\rm Fe}$ distribution broadening in the ACE data and the bimodal $Q_{\\rm Fe}$ signature in the STA data. Therefore, we find aspects of {\\it both} the BM and FC simulation-derived charges states present in the 2007 May 21--23 multispacecraft ICME observations.  Overall, our analysis and results provide the most comprehensive view to date of the global structure of enhanced composition within ICMEs. More investigation is needed to determine whether the structures observed in this case study are common; in particular, the appearance of enhanced charge states surrounding the flux ropes is at odds with previous statistical results indicating that enhanced charge states are more common at the flux rope center. Finally, we show that the ability to derive charge state information from MHD models provides an important and potentially very useful constraint for these models."
    },
    "1208/1208.4947_arXiv.txt": {
        "abstract": "{One of the most challenging steps in planet formation theory is the one leading to the formation of planetesimals of kilometre size. A promising scenario involves the existence of vortices able to concentrate a large amount of dust and grains in their centres. Up to now this scenario has been studied mostly in 2D razor thin disks. A 3D study including, simultaneously, the formation and resulting dust concentration of the vortices with vertical settling, was still missing.} {The Rossby wave instability self-consistently forms 3D vortices, which have the unique quality of presenting a large scale vertical velocity in their centre. Here we aim to study how this newly discovered effect can alter the dynamic evolution of the dust. } {We performed global 3D simulations of the RWI in a radially and vertically stratified disk using the code MPI-AMRVAC. After the growth phase of the instability, the gas and solid phases are modelled by a bi-fluid approach, where the dust is considered as a fluid without pressure. Both the drag force of the gas on the dust and the back-reaction of the dust on the gas are included. Multiple grain sizes from $1mm$ to $5cm$ are used with a constant density distribution.} {We obtain in a short timescale a high concentration of the largest grains in the vortices.  Indeed, in $3$ rotations the dust-to-gas density ratio  grows from $10^{-2}$ to unity leading to a concentration of mass up to that of Mars in one vortex. The presence of the radial drift is also at the origin of a dust pile-up at the radius of the vortices. Lastly, the vertical velocity of the gas in the vortex causes the sedimentation process to be reversed,  the $mm$ size dust is lifted and higher concentrations are obtained in the upper layer than in the mid-plane.} {The Rossby wave instability is a promising mechanism for planetesimal formation, and the results presented here can be of particular interest in the context of future observations of protoplanetary disks.} ",
        "introduction": "In the current formation theory, planets are supposed to be built from colliding planetesimals of kilometre or larger size, but the formation of these planetesimals is still an issue \\citep{CY10}.  Due to their intermediate sizes, they cannot stick through chemical bonds and van der Vaals forces, as opposed to microscopic dust \\citep{D09, BL08}. Moreover, their gravitational fields are too small to retain collision fragments \\citep{B00}. Besides, the gas is partially supported by the radial pressure gradient and is therefore sub-keplerian. The solids in keplerian rotation feel a head-on wind which slows them down. This drag force induces a radial drift toward the central star on timescales as short as hundreds of years for meter size solids \\citep{WEI77}. This timescale appears to be even shorter when compared to the planet formation timescale of a few million years.  Multiple scenarios have been proposed to overcome this difficulty.  The streaming instability \\citep{YG05,JOH07} is an hydrodynamical instability that grows partially thanks to the strong coupling between gas and dust. But its domain of interest only includes regions with an increased dust-to-gas density ratio, compared with the standard value of $\\rho_d/\\rho_g \\sim 10^{-2}$. Another possibility, that does not exclude the previous one, is the presence of vortices in the protoplanetary disk. In this scenario the meter size barrier is outstripped by the presence of vortices that can concentrate solids in their centre and accelerate the growth process. \\citet{BAR95} have shown, with an analytical approach, that anticyclonic vortices effectively concentrate solids in their centre, and this idea was further studied by \\citet{TBD96}. As this concentration effect by the vortices shares the same physics as the radial transport of solids toward the disk centre, the highest concentration is expected for the fastest drifting solids. Whereas these studies ignored the vertical structure of the disk, a 3D approach was proposed by \\citet{SHE06} and \\citet{HK10}. Numerical simulations of these dusty vortices have also been performed in 2D \\citep{BCP99,GL99-2,GL00} and 3D \\citep{JAB04} in order to investigate their ability to concentrate solids.  However these works leave unspecified the formation mechanism for the vortices and their long term evolution. \\citet{KLA03} and \\citet{KLA04} have proposed a non-linear hydrodynamical instability growing in an entropy gradient \\citep{PJS07,PSJ07}, called the baroclinic instability, to form such structures. \\citet{LP10} showed that these vortices are stable structures in 3D, whereas the MHD approach of \\citet{LK11} proved that this instability can form vortices only in the dead-zone. However the vortices migrate radially due to the radial pressure gradient \\citep{PLP10} and their long term stability is not clear. Another proposed formation mechanism for the vortices is the Rossby wave instability (RWI). This is a linear instability that grows in the region of a pressure extremum, which avoids vortex migration \\citep{MKC12}. This instability was first studied in 2D both analytically \\citep{LOV99} and numerically \\citep{LI01,VAR06}. The concentration of solids in Rossby vortices has been explored in numerical simulations \\citep{INB06,LYR09}. Recently, \\citet{RJS11} have proposed the RWI as an explanation for the non-axisymmetric submillimeter images of some transition disks. Moreover, an equivalent of the RWI can also exist in a magnetised disk \\citep{TAV06,YL09}, extending its domain of application. One intriguing characteristic of these vortices is that the vertical displacements of gas in the vortices centres over the whole vertical scale height of the disk. This was first obtained in numerical simulations \\citep{MEH10} and then confirmed analytically \\citep{MYL12,L12}. These structures are of particular interest for the study of the dust concentration. Indeed, not only can they accelerate the concentration in the mid-plane at the centre of the vortices due to the downward flow, they also replenish the upper region of the disk with small particles thanks to the upward flow.   The goal of this paper is to investigate the behaviour of solid particles in 3D vortices formed by the RWI. To this end, we perform full 3D simulations of a disk subject to the RWI. When the vortices are formed, we follow the joint evolution of the gas and the dust, for multiple dust species, through bi-fluid simulations. Our paper is organised as follows. We first present the formation of the vortices, detailing the characteristics of the RWI, and explaining the numerical methods and the results of this gas-only simulation. Section \\ref{sec:dust} deals with the  model we have used for the dust, which is considered as a pressure-less fluid, as well as the limits it sets. The resulting simulations are presented, showing the concentration of the dust in the mid-plane and its vertical stratification. We discuss these results in section \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  We have studied the concentration of dust particles in 3D vortices. To our knowledge this is the first time the dust-trapping mechanism has been explored in stable three-dimensional Rossby vortices. We have first done a simulation of the self-consistent formation the vortices by the Rossby wave instability before including the dust particles. An important difference with the previous studies using analytical vortices (\\emph{e.g.} \\citealt{K81}) is the presence of a vertical velocity in the inner part of the vortices. We have presented the dust-trapping properties of the 3D Rossby vortices. This mechanism is very efficient when the dust is only partially coupled to the gas ($\\Omega_K^0\\tau_s^0=0.5$), and a high dust density is reached in the mid-plane. The estimation of the dust mass concentrated in the vortices gives a value of approximately the mass of Mars in a sphere of radius $0.1AU$ with a higher density reached in the centre.  Those particles more coupled to the gas show a larger density in the upper region of the disk. For these intermediate size grains ($mm$ to $cm$ sizes), there is a competition between sedimentation toward the mid-plane and lifting toward the surface by the vertical velocity of the vortices. This high dust density in the upper region is of particular interest in the context of the forthcoming observation of protoplanetary disks but this needs to be confirmed by the use of a radiative transfer code with a full 3D approach. The results may differ from those obtained with a razor thin disk approach \\citep{WK02,RJS11}.  This mechanism is accordingly more convincing since its efficiency is the highest for the fastest drifting solids, namely when the stopping time is of the order of unity. Future work should study the growth of the instability in the presence of dust to understand how the dust modifies the instability. Moreover a simulation with multiple dust sizes in a multi-fluid simulation is necessary to understand how the small dust is concentrated if the vortices start to be accelerated by the larger particles. Furthermore, we have considered a gaussian pressure bump without considering its formation process, which would give the proper shape of the bump, and then the characteristics, including amplitude, of the RWI. A global study, including accretion processes in the disk, is still needed to give the amplitude of the bump, the consequent number of vortices and then the amount of dust concentrated in such vortices. An important step in this direction was done by \\citet{LM12}.  Finally, in this paper we associated a stopping time with a dust size and fixed density, but the opposite approach can also be used to study the behaviour of dust grains of the same size but different composition."
    },
    "1208/1208.3560_arXiv.txt": {
        "abstract": "Using observations of pulsars from the Parkes Pulsar Timing Array (PPTA) project we develop the first pulsar-based timescale that has a precision comparable to the uncertainties in international atomic timescales.  Our ensemble of pulsars provides an Ensemble Pulsar Scale (EPS) analogous to the free atomic timescale \\'Echelle Atomique Libre (EAL). The EPS can be used to detect fluctuations in atomic timescales and therefore can lead to a new realisation of Terrestrial Time, TT(PPTA11).  We successfully follow features known to affect the frequency of the International Atomic Timescale (TAI) and we find marginally significant differences between TT(PPTA11) and TT(BIPM11).  We discuss the various phenomena that lead to a correlated signal in the pulsar timing residuals and therefore limit the stability of the pulsar timescale. ",
        "introduction": "Atomic frequency standards and clocks are now the basis of terrestrial time keeping. Many countries distribute a local atomic timescale. These are combined by the Bureau International des Poids et Mesures (BIPM) to form International Atomic Time (or Temps Atomique International, TAI) which is published in the form of differences from the national timescales\\footnote{The differences between TAI and various other timescales can be obtained from the ``Circular T\" publication available from \\url{http://www.bipm.org/en/scientific/tai/}.  The difference between TT(BIPM) and TT(TAI) is provided at \\url{ftp://tai.bipm.org/TFG/TT(BIPM)}.}.  TAI is the basis for both Coordinated Universal Time (UTC), used for the dissemination of time signals, and Terrestrial Time (TT).  TT is formed by referencing individual clocks to the Earth's geoid. Throughout this paper, we refer to TT(TAI) as terrestrial time realised by TAI.  Once published, TAI itself is never revised, but the BIPM publishes another realization of TT which is computed every year and labelled TT(BIPMYY), where YY corresponds to the year of the most recent data used. For instance, in this paper we refer to TT(BIPM11) as the most recent post-corrected realisation.  The difference between TT(BIPM11) and TT(TAI) is shown in the top panel of Figure~\\ref{fg:tai_bipm} and clearly shows a drift between the time standards of $\\sim 5$\\,$\\mu$s since 1994.  The stability of TAI is obtained from a large number of atomic clocks whereas the accuracy of TAI is set from a few primary frequency standards (Arias, Panfilo \\& Petit 2011)\\nocite{app11}. Initially, the free atomic timescale \\'Echelle Atomique Libre (EAL) is produced from the weighted average of the timescales of several hundred atomic clocks around the world.  This timescale is not in accord with the second as defined in the International System of Units (SI). Therefore, to form TAI (which does conform to the SI second) from EAL, various frequency adjustments are necessary.  These are determined using primary frequency standards.  Frequency adjustments are generally made slowly, a process referred to as ``steering''.  In 1996, a decision was made to change the realization of the SI second that resulted in a frequency shift of about $2\\times10^{?14}$. That shift was progressively introduced into TAI over a period of two years.  As TAI itself is never retroactively corrected, only the post-corrected versions of TT, e.g., TT(BIPM11), have the earlier data corrected.  This leads to the ``bump'' that we observe in Figure 1 around the year 1998.  Although numerous clocks are used in forming TAI and there is continuous development of atomic clocks, stability over decades is difficult to measure and maintain.  It is therefore desirable to have an independent precise timescale valid on such long intervals.  In this paper, we describe the development of such a timescale based on the rotation of pulsars.  \\begin{figure} \\includegraphics[width=6cm,angle=-90]{tai_bipm.ps} \\caption{The top panel shows the difference between TT(BIPM11) and TT(TAI) since the year 1994.  The bottom panel shows the same, but after a quadratic polynomial has been fitted and removed.} \\label{fg:tai_bipm} \\end{figure}  Radio pulsars are rotating, magnetised neutron stars that radiate beams of electromagnetic waves.  For a fortuitous line of sight to the pulsar, these can be observed at the Earth as pulses.  The pulse times of arrival (ToAs) from the brightest and fastest-spinning pulsars can be measured with a precision of $\\sim100$\\,ns in an observation time of $\\sim$\\,1\\,hour.  This precision is significantly worse than that obtainable from atomic clocks, but, in contrast to individual clocks, can be maintained for a very long time.  We note that a pulsar-based timescale provides: \\begin{itemize} \\item{an independent check on terrestrial timescales using a system that is not terrestrial in origin.} \\item{a timescale based on macroscopic objects of stellar mass instead of being based on atomic clocks that are based on quantum processes.} \\item{a timescale that is continuous and will remain valid far longer than any clock we can construct.} \\end{itemize}  In order to develop a pulsar-based timescale, all phenomena affecting the pulse ToAs must be taken into account.  These are incorporated into a ``pulsar timing model'' that contains the pulsar's astrometric, rotational and orbital parameters, the effects of the interstellar medium and the motion of the Earth about the solar system barycentre. Timing residuals are the difference between the arrival times converted to the solar system barycentre and predictions of those times based upon the timing model (see e.g., Edwards, Hobbs \\& Manchester 2006 for details)\\nocite{ehm06}.  Non-zero residuals can result from an incorrect conversion from the measured ToAs to barycentric arrival times.  Our ability to convert to barycentric arrival times relies, for instance, upon the accuracy of the solar system ephemeris.  Many pulsars also display irregularities in rotation and changes in pulse shape that make timing difficult (e.g., Lyne et al. 2011 and references therein\\nocite{lhk+10}).  A subset of pulsars, the ``millisecond pulsars'', have shorter pulse periods and much more stable rotation than the ``normal pulsars\".  However, precise observations of millisecond pulsars show some unexplained timing irregularities which we refer to as ``timing noise''.  Some of the variations in the timing residuals are caused by processes that are correlated between different pulsars.  These can be identified by observing an ensemble of pulsars, a so-called ``Pulsar Timing Array'' (PTA, e.g., Foster \\& Backer 1990)\\nocite{fb90}.  Errors in the terrestrial time standard will introduce exactly the same signal in the residuals for each pulsar.  In contrast, errors in the planetary ephemeris used in the timing analysis will induce timing residuals which have a dipolar signature on the sky and gravitational waves propagating past the pulsar and the Earth will induce timing residuals with a quadrupolar signature. As shown later in this paper, it is not possible to obtain an unbiased estimate of the time standard errors simply by forming a weighted average of the timing residuals for different pulsars. This is because of the coupling between the timing model for each pulsar and the measurement of the correlated signal as well as the differing data spans for each pulsar.  In this paper we analyse data from the Parkes Pulsar Timing Array (PPTA) project (Manchester et al. 2012)\\nocite{mhb+12} to develop a pulsar-based timescale which we label an Ensemble Pulsar Scale (EPS). This scale has similarities to the free atomic timescale EAL. The frequency of EAL needs to be steered using primary frequency standards to realise a timescale based on the SI second.  Similarly, since the intrinsic pulsar pulse periods and their time derivatives are unknown for the pulsars in a PTA, the EPS is not an absolute timescale and it must be ``steered'' to a reference timescale which conforms to the SI. This is achieved by first forming timing residuals for each pulsar with respect to the reference timescale, TT(TAI) in our case, and subsequently fitting a quadratic polynomial to the residuals. Fluctuations in the reference timescale with respect to the EPS can be identified and used to provide a set of corrections to that realisation of TT, thereby realising a new pulsar-based timescale. We refer to the timescale derived in this paper as TT(PPTA11).  The bottom panel of Figure~\\ref{fg:tai_bipm} shows the difference between TT(BIPM11) and TT(TAI) after a quadratic polynomial has been fitted and removed. It is this signal that we expect to see in comparing TT(PPTA11) with TT(TAI).   Earlier attempts to develop a pulsar timescale have been made by Guinot \\& Petit (1991)\\nocite{gp91}, Petit \\& Tavella (1996)\\nocite{pt96}, Rodin (2008)\\nocite{rod08} and Rodin \\& Chen (2011)\\nocite{rc11}\\footnote{Note that some authors (e.g., Petit \\& Tavella 1996\\nocite{pt96}; Rodin, Kopeikin \\& Ilyasov 1997\\nocite{rki97}) have considered using the orbital parameters of binary pulsars to provide a pulsar-based timescale.}.  We will show below that, in contrast to our method, these earlier attempts did not account fully for the effects of fitting a pulsar timing model. They also have not been applied to high precision observations for a large number of pulsars.  In \\S2 we describe the signal that is potentially measurable using pulsar observations.  \\S3 describes the observations used in this paper.  \\S4 contains details of the method applied. \\S5 presents the application of our method to actual data and contains a discussion on the result. \\S6 summarises the results.  The algorithm presented here has been included in the \\textsc{tempo2} pulsar-timing software package (Hobbs, Edwards \\& Manchester 2006)\\nocite{hem06}.  Usage instructions are given in Appendix A. ",
        "conclusions": "We have developed a new algorithm for determining the correlated signal in the timing residuals for multiple pulsars.   Any errors in the reference timescale will lead to such a correlated signal.  By comparing our measurements of pulse arrival times to TT(TAI) we have confirmed that we can recover the effects of the deliberate steering of TAI. We have not identified any significant discrepancies with TT(BIPM11), but have noted a marginal discrepancy between 1995 and 2003.  Other phenomena, such as an isotropic, stochastic, gravitational wave background will also lead to a correlated signal in pulsar data sets, but we show that such phenomena are not likely to affect our results.  In the future it is likely that pulsar data sets will be processed in a manner that will simultaneously identify: irregularities in the time standard; errors in the solar system ephemeris; and gravitational waves.  By combining observations from numerous telescopes, such future data sets will significantly improve on the pulsar-based timescale that is presented here."
    },
    "1208/1208.5493_arXiv.txt": {
        "abstract": "The solar magnetic activity cycle causes changes in the Sun on timescales that are relevant to human lifetimes. The minimum in solar activity that preceded the current solar cycle (cycle 24) was deeper and quieter than any other recent minimum.  Using data from the Birmingham Solar-Oscillations Network (BiSON), we show that the structure of the solar sub-surface layers during the descending phase of the preceding cycle (cycle 23) was very different from that during cycle 22. This leads us to believe that a detailed examination of the data would have led to the prediction that the cycle-24 minimum would be out of the ordinary. The behavior of the oscillation frequencies allows us to infer that changes in the Sun that affected the oscillation frequencies in cycle 23 were localized mainly to layers above about 0.996R$_\\odot$, depths shallower than about 3000 km. In cycle 22, on the other hand, the changes must have also occurred in the deeper-lying layers. ",
        "introduction": "The solar magnetic activity cycle, usually referred to as the solar cycle, causes changes on the Sun on time scales relevant to human life on earth. The maxima of the cycles are marked by large numbers of sunspots, as measured by the International Sunspot Number (ISN). Also prevalent are solar flares and emissions that are a direct result of areas of strong magnetic fields on the Sun, e.g., radio emission at a wavelength of 10.7 cm (RF10.7). The effect that these flares have on the Earth has led to many efforts to try to predict levels of solar activity led by the Space Weather Prediction Center of the National Oceanic and Atmospheric Administration.\\footnote[3]{http://www.swpc.noaa.gov/SolarCycle/}  The minimum that preceded the current solar cycle (cycle 24) was unusually quiet and long, and defied all existing  predictions. Most solar activity markers, such as the RF10.7, were at an all-time low; the polar magnetic fields were weak and the amount of local flux was small (see e.g. articles in ASPCS vol. 428). Observations of interplanetary scintillation (IPS) showed solar-wind turbulence (a quantity related to solar magnetic fields) had been steadily dropping for a decade before the deep minimum (Janardhan et al. 2011). Such a low minimum had not been seen since detailed solar-cycle observations began with the advent of the satellite era. The predictions for this solar minimum, including the timing for when the Sun would emerge from it, turned out to be incorrect raising questions about the validity of the prediction techniques (Hathaway 2010).  The activity cycle has its origins in processes occurring in the solar interior. While the usual solar-cycle related observations deal with phenomena on the solar surface or beyond, helioseismic data, i.e., data on solar oscillations, provide the means to probe beneath the solar surface. A number of observational programs \uf02d such as the Global Oscillations Network Group (GONG), the Michelson Doppler Imager on board the Solar and Heliospheric Observatory (SoHO) and the Birmingham Solar Oscillations Network (BiSON)  have been collecting data on solar oscillations. Helioseismologists use the frequencies of solar oscillations to derive the structure and dynamics of the solar interior (e.g. Christensen-Dalsgaard 2002).  Solar oscillations respond to the changing solar activity. It has been shown that the frequencies of solar oscillations increase with activity (Woodard \\& Noyes 1985; Elsworth et al. 1990; Libbrecht \\& Woodard 1990). The frequency shifts are predominantly a function of frequency alone, and all the degree dependence can be removed by correcting for mode inertia (mode inertia determines how a mode responds to a given perturbation, high-frequency and high-degree modes have lower inertia than low-frequency or low-degree modes and hence their frequencies are changed more easily). The nature of these solar-cycle related variations is such that most of the changes to the structure of the Sun are believed to be restricted predominantly to a thin sub-surface layer (Nishizawa \\& Shibahashi 1995). Changes to solar structure in the deep interior are small and the inferred change in the sound-speed at the base of the convection zone is only of the order of $10^{-5}$ (Baldner \\& Basu 2008).  Dynamical variations do however penetrate the sub-surface convection zone more deeply, and analysis of GONG and MDI data has shown that a pole-ward flow of material seen in the solar interior at the beginning of cycle 23 was not seen at the beginning of cycle 24 (Howe et al. 2009; Antia \\& Basu 2010). Some theoretical studies have suggested that differences in behavior between solar cycles are caused by variations of flows, particularly the meridional flow, inside the Sun (Nandy et al. 2011).  Neither the GONG nor MDI observations could predict that cycle 23 would end in a deep minimum, since there were no data from prior cycles to compare against. However, BiSON has now been collecting solar oscillations data for almost three complete solar cycles stretching back over 30 years. In this work we show that an early comparison of the oscillation frequencies from the descending phases of cycles 22 and 23 would have revealed that we were headed towards a peculiar minimum. Indeed, as we show below, we would have known that the entire descending phase of cycle 23 was strange compared to that of cycle 22. ",
        "conclusions": "{ The minimum preceding solar cycle 24 was unusual in its depth and duration. What causes such deep minima is still a matter of debate. We have analyzed 30 years of helioseismic data collected by the Birmingham Solar Oscillation Network to compare the solar cycle 23 with cycle 22 in order to determine whether cycle 23 itself showed peculiarities before the onset of the peculiar minimum.  We compared the solar-cycle related frequency shifts with two activity proxies --- the 10.7 cm flux and the international sunspot number --- to find that while the activity proxies matched the frequency shifts during cycle 22, this was not the case during cycle 23. In the high- and medium-frequency ranges the proxies matched the frequency shifts in the ascending phase of the solar cycle, but started showing a marked mismatch just before the cycle 23 maximum and that mismatch continued in the declining phase of the cycle. In the low-frequency band, the proxies did not match the cycle-23 shifts at all. This indicates that the peculiarity in the solar cycle started long before the minimum.  Further analysis of the frequency shifts showed that the high- and mid-frequency modes in cycles 22 and 23 changed in a similar manner, but the low-frequency modes did not change much in cycle 23. Thus the Sun had started showing deviations from ``normal'' activity-related behavior long before the onset of the deep, ``peculiar'' minimum. Helioseismic signatures were already present during cycle 23 which indicated that the Sun was behaving in a way that had not been observed before. With the benefit of hindsight, these signatures could have been used to forecast the unusual minimum.  The nature of the changes in solar frequencies in cycle 22 compared with that of cycle 23 suggests that solar-cycle related changes during cycle 22 occurred deeper in the Sun than in solar cycle 23. Effectively, the layer of magnetic field  had become thinner after cycle 22 and was confined to shallower layers of the Sun.  By using well-established helioseismic techniques we are able to localize these changes. We note also, that the frequency shifts seen in the low-frequency band are not at the low level seen at previous minima. We can perhaps conjecture that there is some trapped flux that is unable to make its way to the surface. In all, there is now the potential for further quantitative investigation into the behavior of the magnetic cycle of the Sun. }"
    },
    "1208/1208.0878_arXiv.txt": {
        "abstract": "Even though the solar wind is highly supersonic, intense ultra-low frequency (ULF) wave activity has been detected in regions just upstream of the bow shocks of magnetized planets. This feature was first observed ahead of the Earth's bow shock, and the corresponding region was called the ULF wave foreshock, which is embedded within the planet's foreshock. The properties as well as the spatial distribution of ULF waves within the Earth's foreshock have been extensively studied over the last three decades and have been explained as a result of plasma instabilities triggered by solar wind ions backstreaming from the bow shock. Since July 2004, the Cassini spacecraft has characterized the Saturnian plasma environment including its upstream region. Since Cassini's Saturn orbit insertion (SOI) in June 2004 through August 2005, we conducted a detailed survey and analysis of observations made by the Vector Helium Magnetometer (VHM). The purpose of the present study is to characterize the properties of waves observed in Saturn's ULF wave foreshock and identify its boundary using single spacecraft techniques. The amplitude of these waves is usually comparable to the mean magnetic field intensity, while their frequencies in the spacecraft frame yields two clearly differentiated types of waves: one with frequencies below the local proton cyclotron frequency ($\\Omega_{\\text{H+}}$) and another with frequencies above $\\Omega_{\\text{H+}}$. All the wave crossings described here, clearly show that these waves are associated to Saturn's foreshock. In particular, the presence of waves is associated with the change in $\\theta_{Bn}$ to quasi-parallel geometries. Our results show the existence of a clear boundary for Saturn's ULF wave foreshock, compatible with $\\theta_{Bn}\\sim45^{\\circ}$ surfaces. ",
        "introduction": "\\label{int}  When the supersonic solar wind plasma from the Sun encounters an obstacle, a bow shock is formed. The incoming solar wind particles (electrons and ions) upstream from the bow shock have no information about the obstacle, except in a region magnetically connected to it. This region is known as the foreshock. At the bow shock, a small fraction of the solar wind particles are reflected in the sunward direction. These backstreaming particles are subjected to the solar wind's \\textbf{E}$\\times$\\textbf{B} drift, where \\textbf{E}$=-$\\textbf{v}$_\\text{sw}\\times$\\textbf{B}$/c$, is the solar wind's convective electric field, \\textbf{B} is the interplanetary magnetic field (IMF), $\\textbf{v}_\\text{sw}$ is the solar wind velocity and $c$ is the speed of light. The \\textbf{E}$\\times$\\textbf{B} drift velocity is the same for all backstreaming particles, and perpendicular to the IMF. As a result, the guiding centers of all backstreaming particles move within the $\\textbf{v}_\\text{sw}$-\\textbf{B} plane, gradually drifting away from the field line tangent to the bow shock toward the inner part of the foreshock and being segregated according to their parallel velocities. Electrons, because of their much smaller inertia, are much less affected by this drift and their presence can be detected right next to the field line tangent to the bow shock. The combination of solar wind and energetic backstreaming electrons results in the production of Langmuir waves at the electron plasma frequency \\citep{F1985,S2004a,S2004b}. Backstreaming ions on the other hand, can drive a number of plasma instabilities \\citep{G1993,CG1997}, leading to the generation of waves. The ion foreshock is then characterized not just by the presence of a small fraction of backstreaming ions, but also by the generation and propagation of plasma waves around the local ion cyclotron frequency.  The understanding of planetary foreshocks is far from complete. The most studied case is the Earth's foreshock, which is reasonably well understood thanks to single and multi-spacecraft measurements \\citep{TR1981,E2005}. The first observations from the Earth's foreshock were made by the dual spacecraft ISEE \\citep{O1977} which identified different types of backstreaming ion distributions: reflected (now called field-aligned beams), intermediate, and diffuse \\citep{G1978,P1981}. The classification of backstreaming ion populations into these three types was made on the basis of two-dimensional velocity distribution functions and energy-time spectrograms. Further results from ISEE demonstrated the existence of gy\\-ro\\-pha\\-se-bun\\-ched and gyrotropic backstreaming ion distributions in the foreshock \\citep{G1983}.  The field-aligned distributions are typically observed at and near the leading edge of the ion foreshock without the presence of waves \\citep{P1979}. Behind the field-aligned beam region, gyrophase-bunched distributions are detected, while diffuse distributions are found even farther away (i.e. downstream) from the ion foreshock boundary. The association of different linear and non-linear waves to different ion distributions was first studied by \\citet{HR1983}. Gy\\-ro\\-pha\\-se-bun\\-ched and diffuse distributions are in fact observed in the presence of ultra-low frequency (ULF) waves. In particular, gyrophase-bunched distributions coexist with ULF quasi-monochromatic waves with substantial amplitudes ($\\delta\\textbf{B}/B\\approx1$) \\citep{M2003}. The production of gyrophase-bunched ions and associated ULF waves has been studied numerically \\citep{HT1985}, theoretically \\citep{M2000} and observationally, using data from WIND spacecraft \\citep{M2001}. On the other hand, non-linear, steepened waves have been found to be associated with diffuse ion distributions \\citep{H1981}.  Evidence of foreshocks has also been found in other planets as well. In the case of Mars, \\citet{T1992} found electrostatic waves which are polarized along the interplanetary magnetic field, and their peak intensity occurs at or near the local solar wind plasma frequency. \\citet{G1987} identified the ULF wave foreshock boundary of Venus through observations made by the PVO (Pioneer Venus Orbiter) magnetometer. Plasma and magnetic field observations from the Voyager 2 spacecraft reveal ULF waves in the solar wind, which are associated with Neptune's, Jupiter's and Uranus' foreshocks \\citep[see for example][]{Be1991,B1987,R1990}. Plasma waves have also been detected upstream from Mercury's bow shock using Mariner 10 measurements \\citep{FB1976}. A brief review on Mercury's foreshock was made by \\citet{Bl2007}.  The region of ULF wave activity is embedded in the ion foreshock, and its boundary is known as the ULF wave foreshock boundary. In the case of the Earth, \\citet{D1976} located this boundary in a statistical way by using combined magnetic field and plasma data from Heos 1. \\citet{GB1986}, using ISSE 1 data introduced the so called solar foreshock coordinates, and determined the position of the ULF wave foreshock boundary. For different IMF cone angles (they split their study into $\\theta_{Bx}=20^{\\circ}$ - $30^{\\circ}$ and $\\theta_{Bx}=40^\\circ$ - $50^\\circ$ data sets), \\citet{GB1986} obtained different orientations of the ULF wave foreshock boundary. Other studies have shown that for $\\theta_{Bx}\\geq45^\\circ$ cone angles, there is a well defined upstream boundary, and this boundary intersects the bow shock at $\\theta_{Bn}\\approx50^\\circ$ \\citep{LR1992a}. Using multi-spacecraft data obtained by Cluster, \\citet{E2005} presented two case studies on the onset of ULF foreshock waves, relating their appearance to changes in the orientation of the IMF. The increased wave activity is associated with the change in $\\theta_{Bn}$ to quasi-parallel geometries ($\\theta_{Bn}<45^\\circ$). These findings are consistent with previous single spacecraft studies.  Although the observations at Earth represent a vital element in the study of the physical processes occurring on planetary foreshocks at large, these phenomena necessarily occur at the particular parameter values relevant for our planet. To explore how these processes change in parameter space, it is just as important to make in situ observations around other planets. For instance, the solar wind properties vary with heliocentric distance. The Parker angle between the stream lines and the radial direction to the Sun at 1 AU is predicted to be about 45$^\\circ$ for a constant solar wind velocity of 429 km$/$s \\citep{TS1980}. At Saturn's distance (approx. 9 AU), the Parker angle is predicted to be about 85$^\\circ$. \\citet{J2008} examined in detail the hourly averaged IMF data provided by Cassini from 13 August 2003 to 14 November 2004, and found an average Parker angle equal to 86.8$^{\\circ}\\pm$0.3$^\\circ$ for a solar wind speed of 500 km$/$s.  The location and shape of planetary bow shocks are determined by the properties of the solar wind flow and by the size and shape of the obstacle to the flow. Saturn's bow shock was first observed by Pioneer 11 in 1979 \\citep{A1980}. The first crossings made by Cassini during Saturn Orbit Insertion (SOI) were analyzed and discussed by \\citet{A2006}. They presented evidence of magnetospheric compression during Cassini's first immersion into the magnetosphere and the properties of the solar wind upstream from Saturn's bow shock for the first six bow shock crossings observed by Cassini. The magnetic signatures of these bow shock crossings showed a clearly defined overshoot and foot regions associated with the quasi-perpendicular geometry ($\\theta_{Bn}>45^\\circ$). Using magnetic field and plasma observations made by Cassini between June 2004 and August 2005, \\citet{M2008} presented a static model of Saturn's bow shock. The model was obtained by fitting a conic section to the first 206 crossings observed by Cassini. Based on observations from Pioneer 11, Voyager 1, Voyager 2 and Cassini \\citet{W2011} derive a small eccentricity for Saturn's bow shock and found variations in the shock subsolar distance associated to variations in the solar wind dynamic pressure.  Using the Magnetospheric Imaging Instrument (MIMI) on board Cassini, \\citet{K2005} showed measurements of the energetic ion population upstream from both the dusk and dawn sides of the Kronian magnetosphere. During the approach phase and first orbits of Cassini, the observations revealed the presence of a series of distinct upstream bursts of energetic hydrogen and oxygen ions up to distances of 120 Saturn's radii. They concluded that these oxygen upstream events must be particles leaking from Saturn's magnetosphere under favorable IMF conditions. However, \\citet{T2007} studied 45 hours of mass-resolved observations of Cassini Plasma Spectrometer (CAPS), which were performed upstream from Saturn's bow shock. The observations show supra-thermal ions composed of H$^+$ and ions with $m/q=2$, presumably solar wind He$^{++}$, with no detectable contribution from magnetospheric water group ions.  \\citet{B1991} reported the first evidence of upstream low-frequency waves in Saturn. During these observations, the spacecraft was magnetically connected to Saturn's bow shock. Their results suggest that these waves were associated to the planet's foreshock. These waves displayed a period of 550 seconds in the spacecraft frame and a relative amplitude of 0.3. Also in the spacecraft frame, the waves are left and right-hand elliptically polarized, and propagate at about 30$^\\circ$ with respect to the ambient magnetic field. During the first three orbits of Cassini spacecraft, \\citet{B2007} present a characterization of low-frequency waves associated with Saturn's foreshock based on Cassini magnetometer \\citep{D2004}. As a result of their survey, they identified two distinct types of waves. They found a large majority of waves with spacecraft-frame frequencies below the local proton cyclotron frequency ($\\Omega_{\\text{H+}}=\\text{eB}_{\\text{0}}/\\text{m}_{\\text{H}}\\text{c}$). These waves are phase-steepened and display a left-hand elliptical polarization as seen by the spacecraft. \\citet{B2007} interpreted these waves as fast magnetosonic waves. This kind of waves is the same presented by \\citet{B1991}. In a second group, they found waves with frequencies above $\\Omega_{\\text{H+}}$, quasi-monochromatic and steepened with a right-hand circular polarization, propagating at small angles with respect to the ambient field. According to \\citet{B2007} these waves could be Alfv{\\'e}n waves similar to those observed at Earth by \\citet{E2003}.  In this paper, we have used Cassini magnetometer data and single-spacecraft techniques to study the morphology of Saturn's ULF wave foreshock. In situ observations made by the Vector Helium Magnetometer (VHM) on board Cassini are presented in section \\ref{obs}. A brief study of ULF waves in the magnetic field associated with Saturn's foreshock is presented in section \\ref{ulf-waves}. In section \\ref{ulf-bound} we show, for the first time, the determination of the outer boundary of Saturn's ULF wave foreshock. For this purpose we identified a large number of crossings of Cassini to (or from) the wave region. The selection criterion for these crossing is described in section \\ref{dtc}. In sections \\ref{res} and \\ref{theta} we introduced the solar foreshock coordinates developed by \\citet{GB1986}, and present our main results. Discussion and our conclusions are summarized in section \\ref{dis} and \\ref{con}. ",
        "conclusions": "\\label{con}  Using Cassini's data from SOI in June 2004 through August 2005, we conducted a detailed survey and analysis of ULF waves upstream from Saturn's bow shock. All the wave events show evidence of magnetic connectivity to Saturn's bow shock, therefore we conclude that these wave events are undoubtedly associated with Saturn's foreshock.  As for their frequencies, we identify two distinct types of wave populations. The most frequently observed have frequencies below $\\Omega_{\\text{H+}}$ and are phase steepened, with periods of the order of 5 to 10 minutes in the spacecraft frame. These waves are often accompanied by precursor whistler wave trains. In agreement with previous works, we suggest that these waves are ion/ion resonant right-hand (fast magnetosonic) mode waves which steepen during the nonlinear regime and emit a dispersive whistler to stop the steepening. We also observe waves with frequencies above $\\Omega_{\\text{H+}}$, which appear as either quasi-monochromatic or steepened, with periods of $\\sim$ 1 minute. The quasi-monochromatic events have a circular right-handed polarization with respect to the mean magnetic field and a propagation slightly oblique with respect to the IMF. A more detailed study will be needed in order to conclude whether the latter are the Saturnian equivalent of the 30 s modes found at Earth. It is worth noticing that we have identified the same two categories than \\citet{B2007} have found, although for a much bigger Cassini MAG's data set.  We identified 21 stationary crossings inbound and outbound from the ULF wave foreshock region. We calculated their solar foreshock coordinates in the $\\textbf{v}_{\\text{sw}}$-$\\textbf{B}$ plane and we have identified for the first time Saturn's ULF wave foreshock boundary. In the $\\mu$-$\\nu$ plane we do not find a clear dependence between the foreshock boundary and the IMF cone angle. We also found that the presence of waves is associated with the change in $\\theta_{Bn}$ to quasi-parallel geometries. Moreover, we find that our determination of the ULF wave foreshock boundary as the surface given by $\\theta_{Bn}=45^\\circ$, is indeed consistent with the linear fit in the $\\mu$-$\\nu$ plane, first proposed by \\citet{GB1986} for the Earth's ULF wave foreshock. In this regard, we speculate that the specular reflection is the candidate process for the reflected ions (gyrophase-bunched and diffuse distributions), since it is the only process to produce a signature in $\\theta_{Bn}$. Finally, we studied the robustness of the $\\mu$-$\\nu$ correlation for a more restrictive criterion. We found that regardless of the criterion that we use, we obtain the same ULF wave foreshock boundary in the $\\mu$-$\\nu$ plane."
    },
    "1208/1208.2085_arXiv.txt": {
        "abstract": "The X-ray transient MAXI~J1836--194 is a newly-identified Galactic black hole binary candidate. As most X-ray transients, it was discovered at the beginning of an X-ray outburst. After the initial canonical X-ray hard state, the outburst evolved into a hard intermediate state and then went back to the hard state. The existing RATAN-600 radio monitoring observations revealed that it was variable on a timescale of days and had a flat or inverted spectrum, consistent with optically thick synchrotron emission, possibly from a self-absorbed jet in the vicinity of the central compact object. We observed the transient in the hard state near the end of the X-ray outburst with the European VLBI Network (EVN) at 5 GHz and the Chinese VLBI Network (CVN) at 2.3 and 8.3 GHz. The 8.3 GHz observations were carried out at a recording rate of 2048~Mbps using the newly-developed Chinese VLBI data acquisition system (CDAS), twice higher than the recording rate used in the other observations. We successfully detected the low-declination source with a high confidence level in both observations. The source was unresolved ($\\leq$0.5~mas), which is in agreement with an AU-scale compact jet. ",
        "introduction": "\\label{sec1}  MAXI~J1836--194 is a new X-ray transient, discovered on 2011 August 30 \\citep{neg11}. The transient was classified as a black hole candidate, because a relativistically broadened iron emission line was revealed by \\emph{Suzaku} observations \\citep{rei12} during its first known X-ray outburst. The X-ray outburst started from the canonical low/hard X-ray state, then it evolved to the hard intermediate state and finally went back to the low/hard state. The outburst was classified by \\citet{fer12} as a ``failed\" outburst because it never entered high/soft X-ray state.  The transient source MAXI~J1836--194 was quite bright in radio and infrared. The Karl~G.~Jansky Very Large Array (VLA) and the RATAN-600 radio observations revealed an optically-thick spectrum with a flux density of 20~--~50 mJy between 5 and 8~GHz \\citep{mil11, tru11}. The RATAN-600 5-GHz light curve indicates that its emission varied frequently by a factor of two on the timescale of days. Very Large Telescope (VLT) observations \\citep{rus11} revealed a secular mid-infrared brightening, reaching 57$\\pm$1~mJy at 12.0~$\\mu$m on 2011 October 11, making it one of the brightest black hole candidate ever detected at 10 microns.  In view of the rare occurrence of such a bright transient source associated with a black hole candidate, we conducted milliarcsecond-resolution VLBI (Very Long Baseline Interferometry) observations of the radio counterpart of MAXI~J1836--194 just after the mid-infrared brightening. In this paper, we report the first VLBI detection of its compact radio core.  \\begin{table*} \\caption{Summary of the VLBI observations.} \\label{tab1} \\scriptsize \\setlength{\\tabcolsep}{6pt}  \\centering \\begin{tabular}{ccccccccc} \\hline Proj. Code & Date      & UT Range       & Stations          & Baseline Lengths & Freq. & Duration & Data Rate       \\\\ & (yyyy-mm-dd) &(hh:mm -- hh:mm)&                & (km)         & (GHz) &  (h) & (Mbps)  \\\\ \\hline CHIN06A  & 2011-10-10   & 08:30 -- 13:00 &     KmShUr     &  1920\\,--\\,3249      & 8.3   &  4   & 2048    \\\\ CHIN06B  & 2011-10-10   & 08:30 -- 13:00 &     KmSh       &       1920           & 2.3   &  4   & 1024    \\\\ RSY01    & 2011-10-17   & 16:10 -- 18:00 & HhWbYsTrOnMcEf &  266\\,--\\,8042       & 5.0   &  2   & 1024    \\\\ \\hline \\end{tabular} \\end{table*}  \\begin{table*} \\caption{The circular Gaussian model fitting results of the detected radio component in MAXI~J1836--194. } \\label{tab2} \\scriptsize \\setlength{\\tabcolsep}{10pt} \\centering \\begin{tabular}{ccccccc}  \\hline Frequency & MJD       & TECOR  &  Right Ascension                        & Declination                       &   Flux Density & Angular Size\\\\ (GHz)     & (day)     &        &                                &                                   & (mJy)          & (mas)       \\\\ \\hline 8.3       & 55844.450 & None   & $18^\\mathrm{h}35^\\mathrm{m}43\\fs44439$ &  $-19\\degr19\\arcmin10\\farcs4825$  &  $6.46\\pm0.33$   & $\\leq0.5$   \\\\ 8.3       & 55844.450 & JPLG   & $18^\\mathrm{h}35^\\mathrm{m}43\\fs44451$ &  $-19\\degr19\\arcmin10\\farcs4869$  &  $6.64\\pm0.32$   & $\\leq0.5$   \\\\ 8.3       & 55844.450 & CODG   & $18^\\mathrm{h}35^\\mathrm{m}43\\fs44452$ &  $-19\\degr19\\arcmin10\\farcs4868$  &  $7.23\\pm0.30$   & $\\leq0.5$   \\\\ \\hline 5.0       & 55851.712 & None   & $18^\\mathrm{h}35^\\mathrm{m}43\\fs44449$ &  $-19\\degr19\\arcmin10\\farcs4917$  &  $4.09\\pm0.56$   & $\\leq0.4$   \\\\ 5.0       & 55851.712 & JPLG   & $18^\\mathrm{h}35^\\mathrm{m}43\\fs44451$ &  $-19\\degr19\\arcmin10\\farcs4923$  &  $4.33\\pm0.45$   & $\\leq0.4$   \\\\ 5.0       & 55851.712 & CODG   & $18^\\mathrm{h}35^\\mathrm{m}43\\fs44455$ &  $-19\\degr19\\arcmin10\\farcs4980$  &  $4.03\\pm0.57$   & $\\leq0.4$   \\\\ \\hline \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\label{sec5} We have presented VLBI observations of the recently detected Galactic variable X-ray source MAXI~1836--194 with the EVN at 5~GHz and the CVN at 2.3/8.4 in October 2011. Despite the unfavourable celestial position for the telescopes located in the Northern hemisphere, this low-declination Southern object has been firmly detected at all three frequencies as an unresolved radio source. The VLBI data allow us to put a firm upper limit on the angular size of the source as 0.5~mas, which corresponds to an AU-scale linear size for a Galactic object and is as expected within the framework of compact synchrotron jet models."
    },
    "1208/1208.5346_arXiv.txt": {
        "abstract": "The \\emph{Planck} 28.5\\,GHz maps were searched for potential Anomalous Microwave Emission (AME) regions on the scale of $\\sim3^{\\circ}$ or smaller, and several new regions of interest were selected.  Ancillary data at both lower and higher frequencies were used to construct spectral energy distributions (SEDs), which seem to confirm an excess consistent with spinning dust models.  Here we present higher resolution observations of two of these new regions with the Arcminute Microkelvin Imager Small Array (AMI SA) between 14 and 18\\,GHz to test for the presence of a compact ($\\sim$10\\,arcmin or smaller) component.  For AME-G107.1+5.2, dominated by the {\\sc Hii} region S140, we find evidence for the characteristic rising spectrum associated with the either the spinning dust mechanism for AME or an ultra/hyper-compact \\textsc{Hii} region across the AMI frequency band, however for AME-G173.6+2.8 we find no evidence for AME on scales of $\\sim 2-10$\\,arcmin. ",
        "introduction": "The \\emph{Planck} satellite (\\cite{2010A&A...520A...1T}, \\cite{2011A&A...536A...1P}) observes the sky in nine frequency bands, covering a range from 30 to 857\\,GHz.  Its wide frequency range potentially allows the detection of AME since the high frequency data above 100\\,GHz can be used to constrain the thermal emission, while the lower frequency data is close to the theoretical peak of the spinning dust emission.  When combined with ancillary data at lower frequencies, the spectra of AME regions can be accurately determined and used to probe their properties on large ($\\sim 1$\\,degreee) scales.  The \\emph{Planck} maps were used to detect several new potential AME regions \\cite{2011A&A...536A..20P} by subtracting a spatial model of known emission mechanisms (synchrotron, free-free and thermal dust) extrapolated from observational or theoretical predictions.  Two, AME-G173.6+2.8 and AME-G107.1+5.2, were selected and ancillary data were used at both higher and lower frequencies to construct SEDs, which contain suggestions of AME consistent with spinning dust emission.  The Arcminute Microkelvin Imager Small Array (AMI-SA) is a radio interferometer situated near Cambridge, UK.  Primarily an SZ survey instrument, the AMI-SA is specifically designed to have high sensitivity to low-surface-brightness emission on scales of $2-$10\\,arcmin.  It operates between 14 and 18\\,GHz, close in frequency to the theoretical peak of the spinning dust emission, the position of which varies between $10-50$\\,GHz depending on grain size composition. The AMI-SA has previously been used both to identify and to characterize spinning dust regions in multiple Galactic (e.g. \\cite{2009MNRAS.394L..46A}, \\cite{2010MNRAS.403L..46S}) and extra-galactic (\\cite{ngc6946}) sources. The higher angular resolution and lower frequency coverage of the AMI-SA make it a highly complementary instrument to the \\emph{Planck} satellite for studies of AME. The synthesised beam of the SA, which is an effective measure of the resolution, is $\\simeq$2\\,arcmin at FWHM, while the \\emph{Planck} maps and ancillary data were smoothed to a common resolution of $\\simeq1^{\\circ}$ for AME detection. ",
        "conclusions": "In the case of both G173.6+2.8 (S235) and G107.1+5.2 (S140) we conclude that the bulk of the excess emission seen by \\emph{Planck} must arise on scales larger than 10\\,arcmin. In the case of G173.6+2.8 we confirm that the dominant source within the \\emph{Planck} aperture, S235, has a spectrum on scales of $2-10$\\,arcmin consistent with optically thin free-free emission. In the case of G107.1+5.2 we demonstrate that the dominant source, S140, has a rising spectrum across the AMI band. This spectrum is consistent with either spinning dust emission or the presence of {\\sc UC/HCHii}. With only the low frequency data available on these scales it is not possible to distinguish between the two mechanisms, however, we conclude that the magnitude of the contribution causing the rising spectral index across the AMI band is not sufficient to explain the measured \\emph{Planck} excess, which must once again arise on scales larger than those probed by AMI. In this case, it is therefore possible that the rising spectrum over the AMI band is also a consequence of spinning dust emission."
    },
    "1208/1208.4164.txt": {
        "abstract": "We demonstrate a newly developed mid-infrared (MIR) planetary nebula (PN) selection technique. It is designed to enable efficient searches for obscured, previously unknown, PN candidates present in the photometric source catalogues of Galactic plane MIR sky surveys. Such selection is now possible via new, sensitive, high-to-medium resolution, MIR satellite surveys such as those from the Spitzer Space Telescope and the all-sky Wide-Field Infrared Survey Explorer (WISE) satellite missions. MIR selection is based on how different colour-colour planes isolate  zones (sometimes overlapping) that are predominately occupied by different astrophysical object types. These techniques depend on the reliability of the available MIR source photometry. In this pilot study we concentrate on MIR point source detections and show that it is dangerous to take the MIR GLIMPSE (Galactic Legacy Infrared Mid-Plane Survey Extraordinaire) photometry from Spitzer for each candidate at face value without  examining the actual MIR image data. About half of our selected sources are spurious detections due to the applied source detection algorithms being affected by complex MIR backgrounds and the de-blending of diffraction spikes around bright MIR point sources into point sources themselves. Nevertheless, once this additional visual diagnostic checking is performed, valuable MIR selected PN candidates are uncovered. Four turned out to have faint, compact, optical counterparts in our H$\\alpha$ survey data missed in previous optical searches. We confirm all of these as true PNe via our follow-up optical spectroscopy. This lends weight to the veracity of our MIR technique. It demonstrates sufficient robustness that high-confidence samples of new Galactic PN candidates can be extracted from these MIR surveys without confirmatory optical spectroscopy and imaging. This is problematic or impossible when the extinction is large. ",
        "introduction": "We present an investigation into the potential of  mid-infrared (MIR) survey data from the Spitzer and WISE space satellite missions as a tool to uncover planetary nebula (PN)  candidates that would be hard or impossible to locate optically. The motivation is to develop MIR PN  candidate selection techniques that can be used to uncover the significant numbers of Galactic PNe which are believed to be hidden behind extensive curtains of dust.  PNe are important astrophysical objects and key windows into late stage stellar evolution. They play a major role in Galactic chemical evolution (Dopita et al. 1997; Karakas et al. 2009), return significant enriched mass to the ISM  (Iben 1995) and are powerful kinematic tracers due to their strong emission lines (eg. Durand, Acker \\& Zijlstra, 1998). Over the last decade Galactic PNe discoveries have entered a golden age due to the advent of narrow-band  Galactic plane surveys of high-sensitivity and resolution. This has been coupled to complementary,  multi-wavelength surveys across near infrared (NIR), MIR and radio regimes in particular from both ground and space-based telescopes. These have provided powerful diagnostic and discovery capabilities (e.g. Cohen et al. 2007, 2011, hereafter Papers 1 and 2; Phillips and Ramos-Larios 2008; Ramos-Larios et al. 2009; Miszalski et al. 2011; Anderson et al. 2012).  The total number of known Galactic PNe  is currently  $\\sim$3000, double what it was a decade ago. This is largely due to the $\\sim$1200 PNe found by the two Macquarie/AAO/Strasbourg H$\\alpha$ PNe surveys (MASH: Parker et al. 2006; Miszalski et al. 2008). MASH PNe were uncovered via scrutiny of the sensitive, arcsecond resolution SuperCOSMOS AAO/UKST H$\\alpha$ survey of the Southern Galactic plane (SHS: Parker et al. 2005). These are now being  supplemented by equivalent discoveries in the Northern Galactic plane (e.g. Mampaso et al. 2006; Viironen et al. 2009a,b;  Sabin et al. 2010) arising from careful searches of the Isaac Newton Telescope Photometric H$\\alpha$ Survey  data (IPHAS: Drew et al. 2005). However, these combined numbers still fall a factor of two short of even the most conservative Galactic PN population estimates (Jacoby et al. 2010) where population synthesis yields 6,600-46,000 PNe depending on whether the binary hypothesis for PN formation is required (e.g. De Marco 2009). Significant numbers of PNe are faint and highly evolved as shown to exist in the local volume sample of Frew \\& Parker (2006) and Frew (2008).  They rapidly become undetectable at distances greater than a few kpc. Such objects currently remain beyond detectability.  However, there are also serious problems with obtaining truly representative samples of PNe across the galaxy due to variable extinction. It is clear that a significant population of Galactic PNe is lurking behind the extensive clouds of gas and dust that obscure large regions in the optical regime. It is the extension of previous PN discovery techniques away from the optically dominant  [OIII] PN emission line in un-reddened spectra to the longer wavelength H$\\alpha$ emission line (that can peer at least partially through the dust), that has led to the major, recent discoveries.  Extension of PN identification techniques  to longer, more favourable wavelengths would clearly be advantageous.  For this pilot study six MIR colour-colour selection criteria were simultaneously applied to the 49 million entries in the GLIMPSE-I point source archive. These criteria are based on sources within three standard errors of the median values of the six unique  MIR colours of the 136 previously known PNe that  fall within the GLIMPSE-I footprint. These are assumed representative of the overall Galactic PN population as given in Paper~1. Only 70 candidate sources were returned. About half turned out to be spurious once the MIR image data were examined. Despite significant extinction, four of the remaining sources had  faint optical detections in the available H$\\alpha$ survey images and are the basis of the optical spectroscopic follow-up. Most Galactic PNe are well resolved  and will not be found in searches of the GLIMPSE-I point-source archive which also has very restricted Galactic latitude coverage. These factors substantially reduce the number of obscured PN candidates returned here.  Section~2 gives the  background importance and context of this MIR PN study. Section~3 briefly describes current knowledge of PN characteristics at non-optical wavelengths and the importance of eliminating mimics. Section~4 describes the candidates' MIR selection. Section~5 gives our new study into MIR false-colour imagery as a powerful diagnostic tool. Section~6 presents the spectroscopic follow-up of the four optical counterparts to our MIR selected PN candidates. They are all confirmed as PNe. An additional, serendipitous PN found adjacent to one of our optically detected MIR sources is also confirmed. Some basic characteristics of the new PNe are also presented. In section~7, we provide some conclusions and suggestions for future work.  % ",
        "conclusions": "We have investigated the potential of the available MIR survey data as a tool to uncover new PN candidates that would be hard or impossible to locate optically. The motivation is to develop robust MIR PN candidate selection techniques that can uncover the significant numbers of Galactic PNe  hidden behind extensive curtains of dust. For this pilot study six MIR colour-colour selection criteria were  applied to the GLIMPSE-I point source archive. These are based on the median values of the unique MIR colours of the 136 previously known PNe that  fall within the GLIMPSE-I footprint (and assumed representative of the overall Galactic PN population). Only 70 candidate sources were returned. Most Galactic PNe are well resolved (e.g. only 5.5\\% of MASH PNe are compact/star-like) and so will not be found in the GLIMPSE-I point-source archive which also has a very restricted Galactic latitude coverage. These factors substantially reduce the number of obscured PN candidates found.  Multi-wavelength image montages of each candidate were examined and four with  faint optical detections in the SHS survey were found.  Spectroscopy  confirmed their likely PN nature. This result represents a clear validation of our general MIR selection technique to identify high quality PN candidates. This is because apart from their faint optical signatures (due to the extinction not being too severe)  they simply fulfil the MIR selection criteria we have developed to identify  PN candidates once their MIR images have been checked. We also confirm the PN nature of PM~1-104, serendipitously uncovered close to one of our MIR selected sources, and update erroneous  positions for both PM~1-104 and  K~3-42 that fall in our sample.  We demonstrate that false-colour images of  MIR selected PN candidates are of high diagnostic value. They enable the environmental context of the MIR point sources to be evaluated, showing that GLIMPSE-I point-source photometry cannot always be taken at face value. We thus rejected 37 (54\\%) of the 70 MIR selected candidates as contaminants due to  adverse MIR background effects, associated dubious photometry or the de-blending of diffraction spikes around bright stars into multiple spurious point-sources. In some cases the character of complementary multi-wavelength optical and NIR data led to rejection. This left  27  high-quality PN candidates not including the four new PN confirmations and the two previously known PNe returned by the search. These results highlight  the dangers of using GLIMPSE-I point source photometry in isolation.  In future we will extend the MIR selection to the GLIMPSE-II and GLIMPSE-3D surveys and expand the search to resolved sources. Sky-background following, application of a threshold a given (low) percentage above this background and then running of pixel-connectivity algorithms will be used to isolate resolved but discrete MIR sources in a process directly  analogous to that used on optical data (e.g. Hambly et al. 2001). This should find resolved MIR sources and yield their integrated MIR magnitudes.  Ultimately we plan to extend our MIR colour-colour techniques to the all-sky coverage of WISE which now enables alternative MIR false-colour images to be constructed. The WISE 3.4$\\mu$m and 4.6$\\mu$m bands are directly equivalent to the first two IRAC bands at 3.6$\\mu$m and $4.5\\mu$m. The final two IRAC bands at 5.8 and 8$\\mu$m  do not have any direct WISE equivalent (with the closest being the WISE 12$\\mu$m band) though the WISE 22$\\mu$m band is similar to the MIPS 24$\\mu$m band. WISE can be used as a substitute for IRAC outside of the GLIMPSE regions with excellent sensitivity but poorer resolution (ranging from six arcseconds for the shorter wavelength bands out to $\\sim$12~arcseconds at 22$\\mu$m).  If we can show that the sensitivity and resolution of the WISE MIR bands can provide the same diagnostic capability as for IRAC, then we can search for MIR PN candidates across the entire sky using essentially the same selection criteria. Examination of known PN detected in WISE (e.g. Fig.~\\ref{PM1-104}) reveals potential in this regard. In this way we can compile MIR-selected PN candidates across the entire area covered by the SHS and IPHAS surveys and also to higher latitudes where there is no H$\\alpha$ coverage. This work is now underway.  %{\\bf Prediction of number of objects expected to be found by this method??}"
    },
    "1208/1208.3675_arXiv.txt": {
        "abstract": "\\noindent Oxygen abundances of 67 dwarf stars in the metallicity range $-1.6<\\feh<-0.4$ are derived from a non-LTE analysis of the 777\\,nm \\oi\\ triplet lines. These stars have precise atmospheric parameters measured by Nissen and Schuster, who find that they separate into three groups based on their kinematics and $\\alpha$-element (Mg, Si, Ca, Ti) abundances: thick-disk, high-$\\alpha$ halo, and low-$\\alpha$ halo. We find the oxygen abundance trends of thick-disk and high-$\\alpha$ halo stars very similar. The low-$\\alpha$ stars show a larger star-to-star scatter in [O/Fe] at a given [Fe/H] and have systematically lower oxygen abundances compared to the other two groups. Thus, we find the behavior of oxygen abundances in these groups of stars similar to that of the $\\alpha$ elements. We use previously published oxygen abundance data of disk and very metal-poor halo stars to present an overall view ($-2.3<\\feh<+0.3$) of oxygen abundance trends of stars in the solar neighborhood. Two field halo dwarf stars stand out in their O and Na abundances. Both G53-41 and G150-40 have very low oxygen and very high sodium abundances, which are key signatures of the abundance anomalies observed in globular cluster (GC) stars. Therefore, they are likely field halo stars born in GCs. If true, we estimate that at least $3\\pm2$\\,\\% of the local field metal-poor star population was born in GCs. ",
        "introduction": "Important clues to understand the formation and evolution of the Milky Way's halo and disk components, as well as any possible connections between them, are imprinted in the photospheric chemical composition of FGK-type dwarf stars. These objects have probably retained the chemical composition of the gas from which they formed, thus being excellent tracers of Galactic chemical evolution (GCE). Moreover, their long lifetimes, in particular for the G and K types, allow us to probe the Milky Way's GCE over many billions of years. Consequently, combined with kinematics and information on stellar ages, chemical composition analyses of FGK dwarf stars have the potential to be useful for reconstructing the history of our Galaxy.  The simplest picture for the Galactic halo formation involves a monolithic collapse which leads to a halo star population showing a strong correlation between orbital eccentricity an overall metal abundance \\citep{eggen62}. Although historically important, this model has long been known to be incomplete. Metallicity determinations of giant stars in globular clusters, for example, led \\cite{searle78} to conclude that the halo was formed in a more ``chaotic'' fashion. Indeed, state-of-the-art simulations show that the properties of the stellar component of galactic halos could be heavily influenced by merging events that occur as the galaxy assembles \\cite[e.g.,][]{abadi03,guo08,read08,stewart08,scannapieco09}. In these scenarios, rather than consisting of a single simple evolving population, the halo is expected to contain sub-structures as remnants of its formation history. The discovery and detailed characterization of halo streams and tidal debris heavily support this idea \\cite[e.g.,][]{majewski93,helmi08,klement10,majewski12}.  Very strong evidence for halo sub-structures in the solar neighborhood has been recently provided in a series of papers by \\cite{nissen10,nissen11}, and \\cite{schuster12}. From a detailed spectroscopic analysis of 94 dwarf stars in the $\\feh$ range from $-1.6$ to $-0.4,$\\footnote{In this work we use the standard definitions: $\\mathrm{[X/Y]}=\\log(N_\\mathrm{X}/N_\\mathrm{Y})-\\log(N_\\mathrm{X}/N_\\mathrm{Y})_\\odot$, and $A_\\mathrm{X}=\\log(N_\\mathrm{X}/N_\\mathrm{H})+12$, where $N_\\mathrm{X}$ is the number density of element X.} \\citet[][hereafter NS10]{nissen10} found that stars with halo kinematics separate into two groups based on their $\\alpha$-element abundances (in their case quantified by the average abundance of Mg, Si, Ca, and Ti). NS10 argue that the halo stars in the low-$\\alpha$ group could have been accreted from satellite galaxies, possibly $\\omega$\\,Cen. Their abundance analysis of heavier elements, particularly Na and Ba/Y, however, showed that the $\\omega$\\,Cen and low-$\\alpha$ halo star connection is weak, unless chemical evolution within the satellite galaxy was different for its inner and outer regions, an idea supported by the observation of an abundance gradient in a dwarf galaxy \\cite[cf.][]{nissen11}. Finally, \\cite{schuster12} show that the low-$\\alpha$ halo stars are about 2--3\\,Gyr younger than the high-$\\alpha$ halo stars and that these two groups exhibit different orbital properties, with the low-$\\alpha$ stars having very eccentric orbits, larger $r_\\mathrm{max}$ (maximum distance from the Galactic center), and larger $z_\\mathrm{max}$ (maximum distance from the Galactic disk).  A very important chemical element missing from the NS10 paper series is oxygen. As the third most abundant element in the universe and in stellar atmospheres, after H and He, and having one of the best identified production sites of all elements as well as reliable supernovae yields, oxygen is crucial for GCE studies. Furthermore, oxygen is a key element in the investigation of abundance variations in globular clusters (GCs). Stars with enhanced Na are known to be depleted in O, i.e., they follow the well-known oxygen-sodium anti-correlation in GCs \\cite[e.g.,][]{gratton04,cohen05,yong05,alves12}. A number of recent studies have investigated the contribution of GCs to the build-up of the field halo population \\cite[][]{yong08,carretta10,martell10,martell11}, but none of them have found field stars with both high Na and low O abundances. Since NS10 have already studied Na in their sample of metal-poor stars, the addition of oxygen allows us to assess to which level the halo field has been contaminated by stars formed in GCs.  Determining reliable oxygen abundances in metal-poor dwarf stars is not a straightforward task. Few spectral features due to oxygen are available in the visible spectrum and all are affected by a number of model uncertainties or severe line blending. Our past experience successfully deriving oxygen abundances from a restricted non-LTE analysis of the 777\\,nm \\oi\\ triplet \\citep[][hereafter R07]{ramirez07} now allows us to infer them. In this work, we derive oxygen abundances for as many as possible of the stars in the NS10 study in order to better understand the nature of the two distinct halo populations in the solar neighborhood. ",
        "conclusions": "Non-LTE oxygen abundances from the 777\\,nm \\oi\\ triplet lines have been derived for as many as possible of the stars in the work by NS10. These authors have derived very precise atmospheric parameters and elemental abundances (excluding oxygen) for their sample stars, allowing them to clearly separate the field halo stars into low-$\\alpha$ and high-$\\alpha$ groups.  We find the run of [O/Fe] abundance ratios with $\\feh$ of high-$\\alpha$ halo and thick-disk stars very similar, while that of low-$\\alpha$ halo stars is systematically lower by about 0.2\\,dex and it has, in general, a larger star-to-star scatter compared to the other two groups. A few additional low-$\\alpha$, low-oxygen abundance halo stars are identified in previously published works. Their kinematic properties strengthen the hypothesis by NS10 that these objects may have originated in dwarf satellite galaxies early in the history of the Milky Way. A connection between the low-$\\alpha$, low oxygen halo stars and $\\omega$\\,Cen is not well established, unless assumptions about the early chemical abundance distribution within this extremely complex globular cluster are made.  Our oxygen abundance data for the three groups of stars studied by NS10 exhibit a behavior that is similar to that of the $\\alpha$-elements. The exceptions are two stars, G53-41 and G150-40, which seem to be the first firm candidates of field halo stars born in globular clusters, although probably not $\\omega$\\,Cen, which has been previously argued as one of the main contributors of low-$\\alpha$ field halo stars. Both G53-41 and G150-40 show the classic signatures of abundance anomalies in globular cluster stars, namely very low oxygen and highly enhanced sodium abundances. Since these properties are seen in 2 of the 67 stars studied in this work, we estimate that the contribution of globular clusters to the local field metal-poor ($-1.6<\\feh<-0.4$) stellar population is at least $3\\pm2$\\,\\%."
    },
    "1208/1208.1736_arXiv.txt": {
        "abstract": "We present an overview of microscopical calculations of the Equation of State (EOS) of neutron matter performed using Quantum Monte Carlo techniques.  We focus to the role of the model of the three-neutron force in the high-density part of the EOS up to a few times the saturation density.  We also discuss the interplay between the symmetry energy and the neutron star mass-radius relation.  The combination of theoretical models of the EOS with recent neutron stars observations permits us to constrain the value of the symmetry energy and its slope. We show that astrophysical observations are starting to provide important insights into the properties of neutron star matter. ",
        "introduction": "The knowledge of the Equation of State (EOS) of pure neutron matter is an important bridge between the symmetry energy, and neutron star properties.  The symmetry energy $E_s$ is the difference of nuclear matter and neutron matter energy, it gives the energy cost of the isospin-asymmetry in the homogeneous nucleonic matter.  In the last few years the study of $E_s$ has received considerable attention (see for example Ref.~\\cite{Tsang:2012} for a recent experimental/theoretical review).  The role of the symmetry energy is essential to understand the mechanism of stability of very-neutron rich nuclei, but it is also related to many phenomena occurring in neutron stars. The stability of matter inside neutron stars is very sensitive to $E_s$ and its first derivative. Around saturation density, neutrons tend to decay to protons through the $\\beta$-decay, and the cooling of neutron stars is strongly connected to the proton/neutron ratio as a function of the density.  This ratio is mainly governed by the behavior of $E_s$ as a function of the density.  The inner crust of neutron stars, where the density is a fraction of nuclear densities, is mostly composed of neutrons surrounding a matter made of extremely-neutron rich nuclei that, depending on the density, may exhibit very different phases and properties. The extremely rich phase diagram of the neutron crustal matter is strongly related to the role of $E_s$.  For example it governs the phase-transition between the crust and the core~\\cite{Newton:2011} and $r$-mode instability~\\cite{Wen:2012,Vidana:2012}.  The study of the EOS is particularly difficult because neutron matter is one of the most strongly-interacting fermionic systems. Neutron matter is often modeled by density functionals. Traditional Skyrme models (see for example Ref. \\cite{Stone:2007} and references therein) and relativistic mean-field models (see for example Refs.~\\cite{Fattoyev:2010,Fattoyev:2012}) are two general classes of density functional theories.  Recently, a new methods based on microscopic nuclear Hamiltonians obtained from chiral effective field theories have been proposed. These nuclear forces are obtained following a systematic expansion in terms of momenta of the relevant degrees of freedom, and fit the nucleon-nucleon scattering data~\\cite{Entem:2003}. The nuclear Hamiltonian is then adjusted using renormalization group techniques to make the calculation perturbative~\\cite{Hebeler:2010}. New induced terms generated trough the renormalization scheme have not yet been included though. These effects, as well as non-perturbative effects of three- and four-body forces, could be important~\\cite{Roth:2011}.  The third class of these calculations uses nuclear potentials, like Argonne and Urbana/Illinois forces, that reproduces two-body scattering and properties of light nuclei with very high precision~\\cite{Wiringa:1995,Pieper:2001}. In the latter case, the interaction has small non-local terms, giving the potentials a hard core. In this case calculations can be performed in a non-perturbative framework, and the strong correlations are solved by using correlated wave functions.  In this paper we present results based on quantum Monte Carlo (QMC) methods. QMC methods have proven to be a very powerful tool to accurately study properties of light nuclei~\\cite{Pudliner:1997,Pieper:2008} and nuclear matter~\\cite{Gandolfi:2010} in a very similar way. They provide the unique technique to date to consistently study nuclear systems of different kind, both inhomogeneous and homogeneous matter, with the same accuracy and using the same Hamiltonians. Other techniques have been used to study nuclear and neutron matter and the symmetry energy based on Brueckner-Hartree-Fock theory (see for example Ref.~\\cite{Vidana:2009} and references therein). ",
        "conclusions": "We have presented a theoretical calculation of the neutron matter EOS using QMC methods. This technique permits to study the ground-state of strongly interacting Fermi systems in a full non-perturbative way. Calculations have been performed using a modern nucleon-nucleon interaction that fit the phase shifts with high accuracy. We have studied the effect of using different microscopical models of three-neutron forces, by quantifying their role in the high-density EOS up to 3$\\rho_0$. By performing simulations using Hamiltonians that give different values of the symmetry energy we conclude that, at present, the uncertainty to the EOS is mainly due to the poor constrain of $E_{sym}$ rather than the model of the three-neutron force. From our calculation we have extracted the relation between $L$ and $E_{sym}$ suggesting that they are quite strongly linearly related.  We also provide new constraints from astrophysical observations. By combining the recent analysis of Steiner {\\emph et al.} with an empirical EOS which form is suggested by QMC simulations, we provided a new constrain to the value of the symmetry energy and its slope at saturation~\\cite{Steiner:2012}. The result is compatible with several experimental measurements~\\cite{Tsang:2012}. We find good agreement for the $M-R$ relation of neutron stars given by QMC prediction and from observations.   \\ack{ The author would like to thank J. Carlson for critical comments on the manuscript.  This work is supported by DOE Grants No. DE-FC02-07ER41457 (UNEDF SciDAC) and No. DE-AC52-06NA25396, and by the LANL LDRD program. Computer time was made available by Los Alamos Open Supercomputing, and by the National Energy Research Scientific Computing Center (NERSC).  }"
    },
    "1208/1208.6320_arXiv.txt": {
        "abstract": "The new-found prevalence of extremely low mass (ELM, $\\Mhe <0.2\\,M_\\odot$) helium white dwarfs (WDs) in tight binaries with more massive WDs has raised our interest in understanding the nature of their mass transfer. Possessing small ($\\Me \\sim 10^{-3}\\,M_\\odot$) but thick hydrogen envelopes, these objects have larger radii than cold WDs and so initiate mass transfer of H-rich material at orbital periods of 6--10 minutes. Building on the original work of D'Antona et al., we confirm the $10^6\\,$yr period of continued inspiral with mass transfer of H-rich matter and highlight that the inspiraling direct-impact double WD binary HM Cancri likely has an ELM WD donor. The ELM WDs have less of a radius expansion under mass loss, thus enabling a larger range of donor masses that can stably transfer matter and become a He mass transferring AM CVn binary.  Even once in the long-lived AM CVn mass transferring stage, these He WDs have larger radii due to their higher entropy from the prolonged H burning stage. ",
        "introduction": "Helium core white dwarfs (WDs) are made from $<2.0\\,M_\\odot$ stars when stellar evolution is truncated before the He core reaches the $\\Mhe \\approx 0.48\\, M_\\odot$ needed for the helium core flash. One formation mechanism is significant mass loss due to stellar winds on the red giant branch (RGB) that strips the H envelope \\citep{ddro96}, typically leading to $\\Mhe=0.4-0.48\\, M_\\odot$ (\\citealt{hansen05,kbh+07}; \\citealt*{ksp07}).  Another mechanism is a common envelope induced by binary interactions (\\citealt{il93}; \\citealt*{mdd95}), making extremely low-mass (ELM) He WDs ($\\Mhe<0.20\\,M_{\\odot}$ or so) when the interaction occurs at the base of the RGB (see \\citealt*{vkbk96}).  These ELM He WDs were first seen as companions to millisecond pulsars \\citep[e.g.,][]{bvkkv06} or in high proper motion catalogs \\citep{kvo+06,kv09}, but the advent of the Sloan Digital Sky Survey \\citep{elh+06} and other surveys revealed many additional ELM WDs \\citep[][and Figure~\\ref{fig:wdmerg}]{ksp07,bmtl09,mbtl09,kbap+10,mgs+11,kvk10,sks+10,kbk+11,kbap+11,pmg+11,marsh11,bkh+11,kbh+11,vtk+11,kbap+12}.  ELM WDs were predicted to possess stably burning H envelopes ($\\Me \\sim 10^{-3}-10^{-2}\\, M_\\odot$) that keep them bright for Gyrs \\citep{sarb02,pach07}, and this has certainly aided the recent detections \\citep{kbap+11,bkapk10}. Though identified by their location in $\\log g-T_{\\rm eff}$ space, few systems have actually had their radii measured with any precision.  Steinfadt et al.'s (\\citeyear{sks+10}) discovery of the eclipsing double WD binary system \\object[NLTT 11748]{\\nltt}\\ \\citep{kv09} allowed for the first geometric measurement of the radius of a ELM WD, finding $R\\approx 0.04\\,R_\\odot$ for the $0.15\\,M_{\\odot}$ He WD, consistent with the presence of a thick stably-burning hydrogen envelope (also see \\citealt*{kvv10}; \\citealt{kapb+10b}).  Additional eclipsing systems \\citep{pmg+11,bkh+11} have led to even more constraints, although for some of the more compact systems it is not clear if the radius is truly the equilibrium radius of the WD or if it has been tidally distorted.  Multiple common envelope phases are possible in the formation scenario for ELM WDs, leading to ELM WDs in binaries with more massive WDs. Many of these will come into contact within 10\\,Gyr (Figure \\ref{fig:wdmerg}).  Indeed, a large number of the known double WD binaries contain ELM WDs ($\\Mhe <0.20 \\, M_\\odot$, circled points) with large ($>0.03 \\, R_\\odot$) radii indicative of a stable H burning shell \\citep{kvo+06,kbap+07,bmtl09,mbtl09,kbap+10,kvk10,mgs+11,bkh+11,kbk+11,vtk+11}. As noted by \\citet{davb+06}, since the time to burn the H envelope can easily exceed the time to reach contact, many of these ELM WDs will come into contact with the remaining H envelope. This raises the possibility for many new phenomena that we begin to explore here.  \\begin{figure} \\plotone{f1.eps} \\caption{\\label{fig:wdmerg} The population of double WDs (triangles and squares, depending on whether one or two members of the binary have radial velocity measurements), sdB/WDs (filled diamonds), and pulsar/WDs (filled stars) with $P_{\\rm orb}< $ day. Binaries to the left of the dashed lines will merge in less than 10 Gyr, 100 Myr, or 1\\,Myr due to gravitational wave losses. Binaries with a ELM He-core WD are circled.  Individual sources of interest are labeled: \\nltt\\ (the first eclipsing system), {CSS~41177} (the second eclipsing system; \\citealt{pmg+11}), {SDSS~J0106$-$1000} and {SDSS~J0651+2844} (two recently discovered short period binaries; \\citealt{kbk+11,bkh+11}), and {KPD~1930+2752} (a binary containing a massive WD and a sdB star with a total mass equal to the Chandrasekhar mass; \\citealt*{mmn00}).  The line at the Chandrasekhar mass $M_{\\rm Ch}=1.4\\,M_{\\odot}$ is where a merger could produce a type Ia supernova in the traditional scenario (but see \\citealt*{vkcj10}). Also see \\citet{kbap+12}.} \\end{figure}  The large radii of the ELM WDs means that Roche lobe overflow (RLO) occurs at larger orbital periods than otherwise expected, so we start in \\S 2 by examining the behavior of the radius of the ELM WD as its H envelope is transferred.  We follow in \\S 3 by outlining the basics of mass transfer and highlighting some of the new possibilities when ELM WDs are donors. There is a more distinct possibility for thermally stable burning of the accreted H and He on the accreting WD and the initial contraction of the ELM WD to mass loss allows for more stable mass transfer than that originally found for cold He WDs \\citep*{mns04}. We perform the full evolution calculations in \\S 4, highlighting the new phases of H mass transfer, the special behavior near the period minimum, and the likelihood that the intruiging object \\object[HM Cnc]{HM Cancri} is one of these systems.  We close in \\S 5 by discussing the implications for AM CVn evolution, and the remaining work needed to resolve the thermonuclear outcomes for the accreted matter. ",
        "conclusions": "We have shown that the unique properties of ELM WDs --- large, nondegenerate H-rich shells supported by stable H burning --- lead to some new phenomena when mass transfer initiates in double-WD binaries. There is a prolonged period of H-rich mass transfer at a low rate during inspiral, with \\hmcnc\\ potentially being the prototype. The change in the mass-radius relation for the donor creates an intrinsically more stable binary that opens up additional phase space for making stable He accreting binaries. This may increase the AM CVn birthrate, potentially alleviating the apparent paucity of progenitor systems.  Those AM CVns which emerged from this progenitor scenario will also have a larger He core radius than expected from an initially cold WD, thereby exhibiting a higher accretion rate at a fixed orbital period than from a cold WD \\citep{db03}.  Prior to the onset of mass transfer, the ELM WDs had Gyrs to undergo diffusive settling and substantial burning of hydrogen.  That clearly allows for the complete sedimentation of the heavier elements from the outermost layers of the WD. Hence, the mass transferred will vary from nearly pure H, to nearly pure He. As we discussed, much work remains to more carefully calculate the thermonuclear outcomes from this mass transfer. If more thermally stable, then these systems may become more observationally detectable due to the higher luminosities. It is also interesting to note the pronounced absence of heavy elements in the x-ray spectra of \\hmcnc\\ \\citep{strohmayer08}, also pointing to an ELM origin that lived a long time prior to mass transfer initiation. If thermally unstable, then the accumulated mass could ignite explosively \\citep{bswn07}, potentially contributing to the increasing number of low-luminosity ``supernovae'' observed locally \\citep[e.g.,][]{kkgy+10}."
    },
    "1208/1208.0506_arXiv.txt": {
        "abstract": "{ The exploration of the relation between galaxy sizes and other physical parameters (luminosity, mass, star formation rate) has provided important clues for understanding galaxy formation, but such exploration has until recently been limited to intermediate redshift objects. } { We use the currently available CANDELS Deep+Wide surveys in the GOODS-South, UDS and EGS fields, complemented by data from the HUDF09 program, to address the relation between size and luminosity at $z\\sim 7$. } { The six different fields used for this study are characterized by a wide combination of depth and areal coverage, well suited for reducing the biases the observed size-magnitude plane. From these fields, we select 153 z-band dropout galaxies. Detailed simulations have been carried out for each of these six fields, inserting simulated galaxies at different magnitudes and half light radius in the two dimensional images for all the HST bands available and recovering them as carried out for the real galaxies. These simulations allow us to derive precisely the completeness as a function of size and magnitude and to quantify measurements errors/biases, under the assumption that the 2-D profile of z=7 galaxies is well represented by an exponential disk function. } { We find in a rather robust way that the half light radius distribution function of $z\\sim 7$ galaxies fainter than $J=26.6$ is peaked at $\\le 0.1$ arcsec (or equivalently 0.5 kpc proper), while at brighter magnitudes high-z galaxies are typically larger than $\\sim$0.15 arcsec. We also find a well defined size-luminosity relation, $Rh\\propto L^{1/2}$. We compute the Luminosity Function in the HUDF and P12HUDF fields, finding large spatial variation on the number density of faint galaxies. Adopting the size distribution and the size-luminosity relation found for faint galaxies at z=7, we derive a mean slope of $-1.7\\pm 0.1$ for the luminosity function of LBGs at this redshift. } { Using this LF, we find that the number of ionizing photons emitted from galaxies at $z \\sim 7$  cannot keep the Universe re-ionized if the IGM is clumpy ($C_{HII}\\ge 3$) and the Lyman continuum escape fraction of high-z LBGs is relatively low ($f_{esc}\\le 0.3$). If these results are confirmed and strengthened by future CANDELS data, in particular by the forthcoming deep observations in GOODS-South and North and the wide field COSMOS, we can put severe limits to the role of galaxies in the reionization of the Universe. } ",
        "introduction": "The advent of the WFC3 instrument onboard HST has opened a new window for the study of galaxy shape, size, and morphology up to very high redshifts (\\cite{windhorst}). The combination of large area, fine resolution, and NIR wavelengths achieved with this powerful instrument allows us to study galaxies in the UV rest frame at $z\\sim 7-8$ with impressive accuracy.  An important parameter that is useful for constraining different models of galaxy formation and evolution is galaxy size, measured through its half light radius (hereafter $Rh$). This quantity gives us an indication of the dynamical state of the galaxy itself and the effects of feedback, minor/major merging, inflows and outflows. In particular, the relation between size and luminosity, or other physical properties (stellar mass, dust extinction, etc.), can give insight into the detailed galactic assembly processes.  With the advent of large surveys, like the Sloan Digital Sky Survey (SDSS, \\cite{sdss}) it has been possible to study the physical properties of local galaxies with great accuracy. Present-day galaxies show a clear correlation between size and stellar mass, with the most massive galaxies having the largest half-light radii (\\cite{shen03,gw08}). This mass-size correlation has been found both for elliptical and spiral galaxies. In particular, \\cite{barden05} pointed out that the same relation holds up to $z\\sim 1$ and its normalization is unchanged with respect to local disk galaxies at least for stellar masses $M\\ge 10^{10}M_{\\odot}$. Recently, \\cite{mosleh12} extended the evolution of the stellar mass-size relation for star-forming galaxies till $z\\sim 7$, finding that the typical size of LBGs increases toward lower redshifts, in agreement with previous measurements at low-z.  Star Forming galaxies at $z\\sim 2-3$, instead, have been extensively studied using ground based spectroscopy, HST imaging and IFU observations. Lyman Break Galaxies (LBGs, \\cite{sh93,madau}) at $z\\sim 2-3$ show a stellar mass-radius relation already established (\\cite{nagy,law}). Similarly, Lyman-$\\alpha$ emitters (LAEs) at the same redshifts present a correlation between their sizes and other physical properties, such as stellar mass, star formation rate (SFR), Spectral Energy Distribution (SED) or dust extinction (\\cite{bond}), with larger galaxies having higher stellar masses, higher dust extinction, and higher SFR. The half light radius of these LAEs at $z=2$, however, is not correlated to the EW in Lyman-$\\alpha$. The stellar mass-radius relation evolves in redshift as $(1+z)^{-1}$, in a manner consistent with the size evolution found by \\cite{bouwens04} and \\cite{ferguson04} for LBGs at $z\\ge 2-5$ and by \\cite{hathi08} at $z\\sim 5-6$.  At $z\\sim 6$, \\cite{sizez6} found a well-defined correlation between measured size and observed magnitudes for 332 photometrically selected LBG candidates: this indicates that a size-luminosity relation could be still in place at high-z. At magnitudes fainter than $z_{\\rm AB}\\sim 28$ there is a clear lack of galaxies larger than 0.2 arcsec, but at such faint levels, the effect of surface brightness dimming is limiting the completeness of large galaxies. At the same redshift, \\cite{Dow2007} found that all the Lyman-$\\alpha$ emitters are more compact than average relative to the observed size-magnitude relation of the large i-dropout sample of \\cite{sizez6}.  The evolution of galaxy size with redshift and luminosity also has important implications for the faint end of the Luminosity Function (LF) and the role of low-luminosity galaxies in the reionization of the Universe. One of the main motivations of this work is to answer to the questions raised in \\cite{grazian11}. In our previous work we have explored different distributions for the half light radius of $z=7$ galaxies. One of the clearest results is that the LF is quite steep ($\\alpha \\sim -2$) if faint galaxies are extended ($Rh\\sim 0.2-0.3 arcsec$ at $J=28-29$) while it turns out to be similar to lower-z LFs ($\\alpha \\sim -1.7$) if objects become smaller at relatively faint magnitudes. This is simply due to the corrections for incompleteness at the faint end of the LF, which are more severe for large and extended galaxies. A steep LF has deep implications for the number of ionizing photons produced by galaxies at $z\\sim 7$: a typical LF with $\\alpha \\sim -1.7$ provides enough light to maintain the reionization process only assuming a large escape fraction of Lyman Continuum photons ($f_{esc}>20\\%$), an IGM that is not clumpy ($C_{HII}<4-6$), and extrapolating this steepness down to very faint flux levels ($M_{1500}=-10$). These constraints are valid under the assumptions of a Salpeter IMF and ignoring the effects of PopIII stars or other exotic sources of ionizing radiation. On the other hand, if faint galaxies are extended, then the resulting LF is quite steep ($\\alpha \\sim -2$) and galaxies alone are able to keep the Universe reionized even for less extreme combinations of escape fraction and clumpiness ($f_{esc}>5\\%$ and $C_{HII}<30$). Thus a detailed analysis on the typical sizes of high-z galaxies and the relation between galaxy size and luminosity is necessary to understand whether galaxies alone are the responsibles for reionization. In \\cite{grazian11} we did not provide a definitive answer to these questions, which therefore motivated the investigation here with a larger sample of galaxies at $z\\sim 7$. Of course, other hypotheses on the sources responsible for reionization at such high-z are possible, like a top-heavy IMF, a large contribution from PopIII stars, or other sources of ionizing photons i.e. high-z AGNs.  Throughout this paper, we will assume a ``concordance'' cosmology with $H_0=70km~s^{-1}~Mpc^{-1}$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. In this cosmological model, an angular dimension of 1 arcsec corresponds to a physical dimension of 5.227  kpc (proper) at z=7. ",
        "conclusions": "Galaxy sizes (half light radii) have been measured for a sample of 153 galaxy candidates at $z\\sim 7$ from the CANDELS HST Multi-Cycle Treasury Program (\\cite{grogin11,koekemoer11}) and HUDF09 project (\\cite{bouwens10c}). In particular, we have used the deep HST imaging database in BVIZYJH bands for the ERS, GDS, P12HUDF and HUDF fields together with the wide area observations in the UDS and EGS fields in the VIJH bands. We select the galaxy candidates at z=7 through the classical z-dropout technique, which has been verified by deep VLT spectroscopy (\\cite{pentericci11}). For the UDS and EGS we use the HST $I_{814}$ band as dropout to select high-z galaxy candidates.  Despite the difficulties of measuring galaxy morphology at $z\\sim 7$, thanks to detailed and extensive simulations, we successfully detect a clear size-luminosity relation for LBGs at high-z. In particular, we found evidences that: \\begin{itemize} \\item {\\bf bright galaxies can be large.} At magnitude brighter than $J=26.6$ (corresponding $\\sim L^\\ast$ at z=7) galaxies have been observed at larger dimensions ($Rh\\sim 0.4 arcsec$ or equivalently 2.3 kpc proper) than at faint magnitudes. Again, Fig.\\ref{ALLmag_rh} shows the extended tail in $Rh$ which is present only for bright galaxies: despite all the deeper fields are sensitive to such extended galaxies, none of them have been found at $J\\ge 26.6$. \\item {\\bf faint galaxies are small.} At $J\\ge 26.6$ the observed sizes of z=7 galaxies are smaller than 0.2 arcsec (corresponding to 1.15 kpc proper). This is evident looking both at Fig.\\ref{ALLmag_rh} and as a result of the detailed simulations summarized by Table \\ref{tab:resu}. \\item {\\bf a size-luminosity relation is already in place at z=7.} The observed dependency of LBG sizes from galaxy luminosity at z=7 has been shown in Fig.\\ref{sizelum}. We have found that a relation exist between these two observables, $Rh\\propto L^\\gamma$, with $\\gamma\\sim 1/3-1/2$. \\end{itemize}  These results have deep implications for our understanding of the reionization of the Universe. The derived size-luminosity relation at $z\\sim 7$ and the fact that faint LBGs have typical half light radii of $\\sim$0.1 arcsec seems to indicate that the slope of the z=7 LF is not extremely steep, due to the correlation between size and $\\alpha$ found in \\cite{grazian11} and recovered here. We detect also a strong field to field variation in the faint regime, with $\\alpha\\sim -1.6$ in the HUDF and $\\alpha\\sim -1.8$ in the P12HUDF field. Using an average value for the two samples, we derived $\\alpha=-1.7\\pm 0.1$. The relevant UV emissivity of LBGs at $z=7$, $\\rho_{UV}$, has been computed by integrating the best fit LF down to $M_{1500}=-10$, and resulted in $\\rho_{UV}=1.1\\cdot 10^{+26}$ $erg~s^{-1}~Hz^{-1}~Mpc^{-3}$. This amount of radiation is not able to keep the Universe re-ionized if the IGM is clumpy ($C_{HII}\\ge 3$) and if the Lyman continuum escape fraction of high-z LBGs is relatively low ($f_{esc}\\le 0.3$). The only configuration that allows a non-neutral Universe due to stellar ionizers in galaxies is if the LBG LF is steeper than $-1.7$, combined with a small clumpiness for the IGM and a high escape fraction of Lyman continuum photons. These are not implausible conditions (see \\cite{Bolton2007} or \\cite{haardt11}), but of course they are extreme assumptions, that could be overcome by simpler explanations, like an additional contribution of faint AGNs (see \\cite{fiore11}), or more exotic explanations (see \\cite{dopita}).  Our results on the size distribution and size-luminosity relation of $z=7$ LBGs have been investigated through detailed and realistic simulations and thus they are robust conclusions. However, the dataset used, especially at the bright side (UDS), is not free from biases due to cosmic variance effects, and we cannot exclude that subtle effects can modify our results. Investigating wide fields searching for $-21\\le M_{UV}\\le -20$ galaxies will be very useful to reinforce our statements. The CANDELS-Wide survey, by observing the COSMOS (\\cite{cosmos}) and the whole EGS (\\cite{egs}) fields, for a total of other 300 sq. arcmin down to J=26.7, will provide $\\sim 70$ additional bright candidates at $z=7$, and would be able to beat down the cosmic variance affecting the bright side of $z=7$ distribution. In addition, WFC3 imaging of $Y\\sim 25$ z-dropout galaxies found with large area ground based imaging (\\cite{ouchi,castellano09,castellano10,bowler12}) will provide useful information on the bright side of the size-luminosity relation not yet covered by the present observations.  Moreover, the CANDELS Deep survey plus the HUDF ultradeep fields (both those already observed and the HST program recently approved for Cycle 19) will extend these results and confirm in a more accurate statistical evidence the trend of the size-luminosity relation at $z\\sim 7$. In particular, the CANDELS-Deep region on the two GOODS fields, covering 150 sq. arcmin down to $J\\sim 28$, will open a very interesting window on the exact shape of the size distribution down to very large half light radii. In addition, the combination of depth and area guaranteed by the CANDELS-Deep survey will decrease the uncertainties on the faint side determination of the $z\\sim 7$ LF due to cosmic variance.  Clearly, the size-luminosity relation found here is the simplest correlation between physical properties of galaxy at $z\\sim 7$. With a full spectroscopic sample in hand, coupled with the deep multi-wavelength dataset available for the CANDELS survey, it will be possible to explore the dependencies of the half light radius with the galaxy stellar mass, SFR, dust extinction or the EW of Ly-$\\alpha$, as already done for star-forming galaxies at smaller redshifts."
    },
    "1208/1208.3196.txt": {
        "abstract": " ",
        "introduction": "Systematic studies of star clusters help us to understand the Galactic structure and star formation processes as well as stellar evolution. By utilizing colour-magnitude diagrams of the stars observed in the optical/near-infrared (NIR) bands, it is possible to determine the underlying properties of the clusters such as age, metallicity and distance. Colour magnitude diagrams of star clusters can be used as a good distance indicators. The distance to a star can be evaluated by trigonometric or photometric parallaxes. Trigonometric parallaxes are only available for nearby stars where {\\em Hipparcos} \\citep{ESA97} is the main supplier for the data. For stars at large distances, the use of photometric parallaxes is unavoidable. In other words the study of the Galactic structure is strictly tied to precise determination of absolute magnitudes.  Different methods can be used for absolute magnitude determination where most of them are devoted to dwarfs. The method used in the Str\\\"omgren's $uvby$-$\\beta$ \\citep{NS91} and in the UBV \\citep{Laird88} photometry depends on the absolute magnitude offset from a standard main-sequence. In recent years the derivation of absolute magnitudes has been carried out by means of colour-absolute magnitude diagrams of some specific clusters whose metal abundances are generally adopted as the mean metal abundance of a Galactic population, such as thin, thick discs and halo. The studies of \\cite{Phleps00} and \\cite{Chen01} can be given as examples. A slightly different approach is that of \\cite{Siegel02} where two relations, one for stars with solar-like abundances and another one for metal-poor stars were derived between $M_{R}$ and the colour index $R-I$, where $M_{R}$ is the absolute magnitude in the $R$ filter of Johnson system. For a star of given metallicity and colour, absolute magnitude can be estimated by {\\em linear} interpolation of {\\em two} ridgelines and by means of {\\em linear} extrapolation beyond the metal-poor ridgeline.  The most recent procedure  used for absolute magnitude determination consists of finding the most likely values of the stellar parameters, given the measured atmospheric ones, and the time spent by a star in each region of the H-R diagram. In practice, researchers select the subset of isochrones with $[M/H]\\pm \\Delta_{[M/H]}$, where $\\Delta_{[M/H]}$ is the estimated error on the metallicity, for each set of derived $T_{eff}$, $\\log g$ and $[M/H]$. Then a Gaussian weight is associated to each point of the selected isochrones, which depends on the measured atmospheric parameters and the considered errors. This criterion allows the algorithm to select only the points whose values are closed by the pipeline. For details of this procedure we cite the works of \\cite{Breddels10} and \\cite{Zwitter10}. This procedure is based on many parameters. Hence it provides absolute magnitudes with high accuracy. Also it can be applied to both dwarf and giant stars simultaneously.  In \\citet{Karaali03}, we presented a procedure for the photometric parallax estimation of dwarf stars which depends on the absolute magnitude offset from the main-sequence of the Hyades cluster. \\cite{Bilir08} obtained the absolute magnitude calibrations of the thin disc main-sequence stars in the optical $M_V$ and in the near-infrared $M_{J}$ bands using the recent reduced {\\em Hipparcos} astrometric data \\citep{Leeuwen07}. \\cite{Bilir09} derived a new luminosity colour relation based on trigonometric parallaxes for the thin disc main-sequence stars with Sloan Digital Sky Survey (SDSS) photometry. \\citet{Yaz10} obtained transformation between optical and near-infrared (NIR) bands for red giants. \\citet{Bilir12} extended this study to middle-infrared (MIR) bands by using Radial Velocity Experiment (RAVE) Third Data Release (DR3) data \\citep{Siebert11}. Both works provide absolute magnitudes for a given photometry from another one.   %TABLE 1 \\begin{table} \\setlength{\\tabcolsep}{2pt} \\center {\\small \\caption{Data for five clusters. We used the data in the first line for each cluster for absolute magnitude calibration, whereas the ones in the second and third lines are for comparison purpose. $l$ and $b$ are the Galactic longitude and latitude of the clusters, the symbol $\\mu_{0}$ indicates the true distance modulus of the cluster.} \\begin{tabular}{ccccccc} \\hline Cluster & $ l$ & $b$ & $E(B-V)$ & $\\mu_{0}$ & $[Fe/H]$ & Ref.\\\\ &  $^{(o)}$  & $^{(o)}$ & (mag)    & (mag)     & (dex)& \\\\ \\hline M92     &  68.34 & +34.86  & 0.025 & 14.72 & -2.15 & (1) \\\\ &        &         & 0.020 & 14.59 & -2.31 & (2) \\\\ &        &         & 0.023 & 14.55 & -2.40 & (3) \\\\ \\hline M13     &  59.01 & +40.91  & 0.020 & 14.38 & -1.41 & (1) \\\\ &        &         & 0.020 & 14.27 & -1.53 & (2) \\\\ &        &         & 0.016 & 14.35 & -1.60 & (3) \\\\ \\hline M71     &  56.75 &$-$ 4.56 & 0.280 & 12.83 & -0.78 & (4) \\\\ &        &         & 0.250 & 13.03 & -0.78 & (2) \\\\ &        &         & 0.220 & 13.10 & -0.80 & (3) \\\\ \\hline M67     & 215.70 & +31.90  & 0.038 &  9.53 & -0.04 & (1) \\\\ &        &         & 0.041 &  9.59 & -0.009& (5) \\\\ &        &         & 0.050 &  9.43 & -0.09 & (6) \\\\ \\hline NGC 6791   &  69.66 & +10.90  & 0.150 & 13.10 &  0.37 & (3) \\\\ &        &         & 0.150 & 13.14 &  0.45 & (7) \\\\ &        &         & 0.100 & 12.94 &  0.37 & (8) \\\\ \\hline \\end{tabular}\\\\ (1) \\cite{Gratton97}, (2) \\cite{Harris10}, (3) \\cite{Brasseur10}, (4) \\cite{Hodder92}, (5) \\cite{Sarajedini09}, (6) \\cite{Hog98}, (7) \\cite{Anthony-Twarog07}, (8) \\cite{Sandage03}. } \\end{table}  %TABLE 2 \\begin{table*} \\centering %\\tiny{ %\\scriptsize{ \\small{ \\caption{ Comparison of the selective and total absorptions evaluated by using different extinction and colour excess ratios. The columns give: (1) the cluster, (2) adopted $E(B-V)$ colour excess, (3) $E(V-J)_p$ the colour excess evaluated by the equation $E(V-J)/E(B-V)=2.25$ and used in the paper, (4)$E(V-J)_c$ the colour excess evaluated by the equation $E(V-J)/E(B-V)=2.30$ for comparison purpose, (5) $\\Delta E(V-J)$ the difference between the colour excesses in columns (3) and (4), (6) $(A_J)_p$  the total absorption evaluated by the equation $A_J/E(B-V)=0.87$ and used in the paper, (7) $(A_J)_c$ the total absorption evaluated by the equation $A_J/E(B-V)=1.70 $, and (8) $\\Delta A_J$ the difference between the total absorptions in columns (6) and (7).} \\begin{tabular}{cccccccc} \\hline (1) &          (2) &          (3) &          (4) &          (5) &          (6) &          (7) & (8) \\\\ \\hline Cluster &   $E(B-V)$ & $E(V-J)_p$ &  $E(V-J)_c$&$\\Delta E(V-J)$ &$(A_J)_p$ &  $(A_J)_c$ &   $\\Delta A_J$ \\\\ \\hline M92 &      0.025 &      0.056 &      0.058 &      0.002 &      0.022 &      0.043 &      0.021 \\\\ M13 &      0.020 &      0.045 &      0.046 &      0.001 &      0.017 &      0.034 &      0.017 \\\\ M71 &      0.280 &      0.630 &      0.644 &      0.014 &      0.244 &      0.476 &      0.232 \\\\ M67 &      0.038 &      0.086 &      0.087 &      0.001 &      0.033 &      0.065 &      0.032 \\\\ NGC 6791 &      0.150 &      0.338 &      0.345 &      0.007 &      0.131 &      0.255 &      0.125 \\\\ \\hline \\end{tabular} } \\end{table*}  In \\citet[][Paper I; Paper II]{Karaali12a, Karaali12b}, we used a procedure for the absolute magnitude estimation of red giants by using the $V_{0}\\times(B-V)_{0}$ and $g_{0}\\times (g-r)_{0}$ apparent magnitude-colour diagrams of Galactic clusters with different metallicities. Here, we extend our procedure to Two Micron All Sky Survey \\citep[2MASS;][]{Skrutskie06} photometry. We aim to estimate $M_{J}$ and $M_{K_s}$ absolute magnitudes for red giants with $J_{0} \\times (V-J)_{0}$ and $K_{{s}_0} \\times (V-K_{s})_{0}$ colour-magnitude diagrams. The outline of the paper is as follows. We present the data in Section 2. The procedure used for calibration is given in Section 3, and Section 4 is devoted to summary and discussion. ",
        "conclusions": "We calibrated the absolute magnitudes $M_{J}$ and $M_{K_{s}}$ for red giants in terms of metallicity by means of the colour magnitude diagrams of the clusters M92, M13, M71, M67, and NGC 6791 with different metallicities. The $J \\times (V-J)$ and $K_{s} \\times (V-K_{s})$ sequences used for the calibration of $M_J$ and $M_{K_{s}}$ are provided from different sources and by different procedures, as explained in the following. The main source is the paper of \\cite{Brasseur10}. The $J \\times (V-J)$ and $K_{s} \\times (V-K_{s})$ sequences for the clusters M92, M13 and  M71 are taken from the tables in \\cite{Brasseur10}, whereas the $J_{0} \\times (V-J)_{0}$ sequence for M67 and  NGC 6791 are obtained by transformation of $V$, $B-V$, $V-I$ data in \\cite{Montgomery93} and by means of the $M_{V} \\times (V-J)_{0}$ diagram in \\cite{Brasseur10}, respectively. Also, the $K_{s} \\times (V-K_{s})$ sequence for M67 are transformed from $V$, $B-V$, $V-I$ data in \\cite{Montgomery93}. The fiducial sequence for NGC 6791 is given in $K_{s} \\times (J-K_{s})$ in \\cite{Brasseur10}. We transformed the $J_{0}$ magnitudes to the $V_{0}$ ones obtained from the $M_V \\times(V-J)_{0}$ diagram and altered the $K_{s} \\times (J-K_{s})$ data to $K_{s} \\times (V-K_{s}$) ones. Thus, we obtained two sets of data for two absolute magnitude calibration, i.e. $J_{0} \\times (V-J)_{0}$  and $K_{{s}_{0}} \\times (V-K_{s})_{0}$ for $M_{J}$ and $M_{K_s}$, respectively. We combined each set of data for each cluster with their true distance modulus and evaluated two sets of absolute magnitudes for the $(V-J)_{0}$ and $(V-K_{s})_{0}$ ranges of each cluster. Then, we fitted $M_{J}$ and $M_{K_s}$ absolute magnitudes in terms of iron metallicity, $[Fe/H]$, by quadratic polynomials, for a given $(V-J)_{0}$ and $(V-K_{s})_{0}$ colour index, respectively. The calibrations cover large ranges, i.e. $1.30 \\leq (V-J)_{0}\\leq 2.80$ and $1.75 \\leq (V-K_{s})_{0}\\leq 3.80$ mag for $M_{J}$ and $M_{K_s}$, respectively.  We evaluated the $M_J$ absolute magnitudes of the clusters M5 ($[Fe/H]=-1.17$ dex) and M68 ($[Fe/H]=-2.01$ dex) by the procedure presented in our study for a set of $(V-J)_{0}$ colour index and compared them with the ones estimated via combination of the fiducial $J_{0}\\times(V-J)_{0}$ sequence and the true distance modulus for each cluster. The total of the residuals lie between -0.08 and +0.34 mag, and the range of 94$\\%$ of them is $0<\\Delta M_J \\leq 0.3$ mag. The mean and the standard deviation of (all) the residuals are $<\\Delta M_J>=0.137$ and $\\sigma_{M_{J}}= 0.080$ mag. For the evaluation of the $M_{K_s}$ absolute magnitudes, we applied the corresponding procedure to the clusters NGC 188 ($[Fe/H]=-0.01$ dex) and M68 ($[Fe/H]=-2.01$ dex). Here again, the range of the residuals, their mean and standard deviation are small, i.e. $-0.10<\\Delta M_{K_s} \\leq+0.27$ mag, $<\\Delta M_{K_s}>=0.109$ and $\\sigma_{M_{K_s}}=0.123$ mag.  We compared the statistical results obtained in this study with the ones in Paper I and Paper II. Table 21 shows that $M_V$, $M_g$, $M_{J}$, and $M_{K_s}$ absolute magnitudes can be estimated with an error less than 0.3 mag. However, one can notice an improvement on $M_{J}$ and $M_{K_s}$ with respect to $M_V$ and $M_g$. The main difference between the data of the clusters in three studies is the large domain of the clusters in $(V-J)_{0}$ and $(V-K_{s})_{0}$ which probably contributed to more accurate calibrations of the apparent $J_{0}$ and $K_{s}$ magnitudes in terms of the corresponding colours with respect to $(B-V)_{0}$ and $(g-r)_{0}$ ones. Accurate calibration in apparent magnitude provided accurate absolute magnitudes. The magnitudes and colours for the cluster M67 used in calibration of $M_{J}$ and $M_{K_s}$  absolute magnitudes are not original, but they are transformed from the $V$, $B-V$, and $V-I$ data by means of the equations of \\cite{Yaz10}. The same case holds for the clusters NGC 188 and M68 which are used in the application of the procedure. Calibrations with high correlation coefficients and small residuals confirm also the equations of \\cite{Yaz10}. As claimed in Paper I and Paper II, there was an improvement on the results therein respect to the ones of \\cite{Hog98}. Hence, the same improvement holds for this study. We quote also the work of \\cite{Ljunggren66}.  %TABLE 21 \\begin{table} \\setlength{\\tabcolsep}{3pt} \\center %\\tiny{ \\scriptsize{ \\caption{Comparison of the results in three studies. The word ``all'' indicates to all the residuals. A subset of the residuals is denoted by a percentage, such as 91\\% or 94\\%.} \\begin{tabular}{cccc} \\hline $\\Delta M$ \u0096 range  & $< \\Delta M >$ & $\\sigma$    & Study \\\\ \\hline $(-0.61, +0.66)$ (all) & 0.05 (91$\\%$) & 0.190 (91$\\%$) & Paper I\\\\ $[-0.40, 0.40]$ (91$\\%$)  &       &       &\\\\ \\hline $(-0.28, +0.43]$ (all) & 0.169 (all) & 0.140 (all) & Paper II\\\\ $( 0.10, 0.40]$, (94$\\%$) &       &       &  \\\\ \\hline $(-0.08, +0.34]$ (all) & 0.137 (all) & 0.080 (all) & This study, $\\Delta M_J$\\\\ $( 0.00, 0.30]$ 94$\\%$ &       &       &  \\\\ \\hline $(-0.10, +0.27]$  & 0.109 & 0.123 & This study, $\\Delta M_{K_{s}}$\\\\ \\hline \\end{tabular} \\label{tab:addlabel} } \\end{table}  Although age plays an important role in the trend of the fiducial sequence of the RGB, we have not used it as a parameter in the calibration of absolute magnitude. A quadratic calibration in terms of (only) metallicity provides absolute magnitudes with high accuracy. Another problem may originate from the red clump (RC) stars. These stars lie very close to the RGB but they present a completely different group of stars. Tables 16 and 19, and Figures 7 and 8 summarize how reliable are our absolute magnitudes. If age and possibly the mix with RC stars would affect our results this should up. Additionally, we should add that the fiducial sequences used in our study were properly selected as RGB. However, the researchers should identify and exclude the RC stars when they apply our calibrations to the field stars.  The accuracy of the estimated absolute magnitudes depends mainly on the accuracy of metallicity. We altered the metallicity by $[Fe/H]+\\Delta[Fe/H]$ in evaluation of the absolute magnitudes by the procedure presented in our study and we checked its effect on the absolute magnitude. We adopted $[Fe/H]=$-2.01, -1.117, -0.01 dex and $\\Delta [Fe/H]=$0.05, 0.10, 0.15, 0.20 dex and re-evaluated the absolute magnitudes for ten $(V-J)_{0}$ and nine $(V-K_{s})$ colour indices for this purpose. The differences between the absolute magnitudes evaluated in this way and the corresponding ones evaluated without $\\Delta [Fe/H]$ increments are given in Table 22. The maximum absolute magnitude differences corresponding to $\\Delta[Fe/H]=0.20$ dex in $M_{J}$ and $M_{K_s}$ lie in the intervals 0.09 $\\leq \\Delta M_{J}\\leq$ 0.31 and 0.05 $\\leq \\Delta M_{K_s}\\leq$ 0.14, respectively, for the metallicities $[Fe/H]=-1.117$ and $[Fe/H]=-2.01$ dex. Whereas, they are about 0.5 mag for the metallicity $[Fe/H]=-0.01$ dex. The mean error in metallicity for 42 globular and 33 open clusters in the catalogue of \\cite{Santos04} is $\\sigma=0.19$ dex. If we assume the same error for the field stars, the probable error in $M_{J}$ and $M_{K_s}$ would be less than 0.3 mag for relatively metal poor stars. Whereas, for the solar metallicity stars, the metallicity error should be $\\sigma_{[Fe/H]}<0.15$ dex in order to estimate accurate absolute magnitudes. That is, the solar metallicities should be determined more preciously.  %TABLE 22 \\begin{table*} \\centering %\\tiny{ \\scriptsize{ \\caption{Absolute magnitudes estimated by altering the metallicity as $[Fe/H]+\\Delta [Fe/H]$. The numerical values of $[Fe/H]$ are indicated in the last column. The absolute magnitudes in column (1) are the original ones taken from Table 16 and Table 19, whereas those in the columns (2)-(5) correspond to the increments 0.05, 0.10, 0.15, and 0.20 dex. The differences between the original absolute magnitudes and those evaluated by means of the increments are given in columns (6)-(9).} \\begin{tabular}{c|ccccc|cccc|r} \\hline \\multicolumn{1}{c}{} & \\multicolumn{5}{c}{$M_{J}$}              & \\multicolumn{4}{c}{$\\Delta M_J$}   &  \\\\ \\hline \\multicolumn{1}{c}{$(V-J)_{0}$} & (1)  & (2)   & (3)  & (4)     & \\multicolumn{1}{c}{(5)} & (6)  & (7)  & (8)  & \\multicolumn{1}{c}{(9)} & \\\\ \\hline 1.50  & -0.200 & -0.139 & -0.076 & -0.012 & 0.055 & 0.061 & 0.124 & 0.189 & 0.255 &  \\\\ 1.70  & -1.775 & -1.703 & -1.629 & -1.550 & -1.469 & 0.072 & 0.147 & 0.225 & 0.306 &  \\\\ 1.90  & -2.773 & -2.707 & -2.639 & -2.568 & -2.496 & 0.066 & 0.134 & 0.204 & 0.277 & $[Fe/H]=-1.117+\\Delta [Fe/H]$\\\\ 2.10  & -3.496 & -3.431 & -3.365 & -3.298 & -3.229 & 0.064 & 0.130 & 0.198 & 0.267 &  \\\\ 2.30  & -4.075 & -4.007 & -3.938 & -3.868 & -3.798 & 0.068 & 0.137 & 0.207 & 0.277 &  \\\\ 2.50  & -4.747 & -4.707 & -4.664 & -4.619 & -4.570 & 0.040 & 0.082 & 0.128 & 0.177 &  \\\\ \\hline 1.50  & -0.986 & -0.952 & -0.917 & -0.879 & -0.841 & 0.034 & 0.070 & 0.107 & 0.146 &  \\\\ 1.70  & -2.498 & -2.481 & -2.460 & -2.436 & -2.408 & 0.018 & 0.038 & 0.062 & 0.090 & $[Fe/H]= -2.01+\\Delta [Fe/H]$\\\\ 1.90  & -3.535 & -3.508 & -3.478 & -3.446 & -3.412 & 0.027 & 0.057 & 0.088 & 0.123 &  \\\\ 2.10  & -4.352 & -4.313 & -4.272 & -4.230 & -4.187 & 0.039 & 0.079 & 0.121 & 0.165 &  \\\\ \\hline \\multicolumn{1}{c}{} & \\multicolumn{5}{c}{$M_{K_s}$}      & \\multicolumn{4}{c}{$\\Delta M_{K_{s}}$}        & \\multicolumn{1}{c}{} \\\\ \\hline \\multicolumn{1}{c}{$(V-K_{s})_{0}$} & (1)   & (2) & (3) & (4) & \\multicolumn{1}{c}{(5)} & (6)     & (7)     & (8)     & \\multicolumn{1}{c}{(9)} & \\multicolumn{1}{c}{}\\\\ \\hline 2.45  & -0.634 & -0.489 & -0.341 & -0.190 & -0.035 & 0.145 & 0.293 & 0.445 & 0.599 &  \\\\ 2.50  & -0.908 & -0.772 & -0.633 & -0.491 & -0.346 & 0.136 & 0.275 & 0.417 & 0.562 & \\multicolumn{1}{c}{$[Fe/H] = -0.01+\\Delta [Fe/H]$} \\\\ 2.55  & -1.159 & -1.031 & -0.900 & -0.766 & -0.630 & 0.128 & 0.259 & 0.393 & 0.529 &  \\\\ 2.60  & -1.390 & -1.268 & -1.144 & -1.018 & -0.889 & 0.122 & 0.245 & 0.372 & 0.500 &  \\\\ \\hline 1.85  & 0.002 & 0.024 & 0.048 & 0.074 & 0.103 & 0.022 & 0.046 & 0.072 & 0.101 &  \\\\ 2.05  & -1.575 & -1.557 & -1.536 & -1.513 & -1.487 & 0.018 & 0.039 & 0.063 & 0.089 &  \\\\ 2.25  & -2.772 & -2.764 & -2.753 & -2.738 & -2.719 & 0.008 & 0.019 & 0.034 & 0.053 & \\multicolumn{1}{c}{$[Fe/H] = -2.01+\\Delta [Fe/H]$} \\\\ 2.45  & -3.707 & -3.695 & -3.679 & -3.661 & -3.639 & 0.012 & 0.027 & 0.046 & 0.068 &  \\\\ 2.65  & -4.499 & -4.468 & -4.434 & -4.398 & -4.361 & 0.031 & 0.065 & 0.101 & 0.138 & \\multicolumn{1}{c}{} \\\\ \\hline \\end{tabular} } \\end{table*}  The absolute magnitudes can be calibrated as a function of ultraviolet excess instead of metallicity, in general. However, an ultraviolet band is not defined in 2MASS photometry. Hence, we calibrated the $M_{J}$ and $M_{K_s}$ absolute magnitudes in term of metallicity which can be determined by means of atmospheric model parameters. Age is a secondary parameter for the old clusters and does not influence much the position of their RGB. The youngest cluster in our study is M67 with age 4 Gyr (Paper I). However, the field stars may be younger. We should remind that the derived relations are applicable to stars older than 4 Gyr.  We conclude that the two absolute magnitudes, $M_{J}$ and $M_{K_s}$, in 2MASS photometry can be estimated for the red giants in terms of metallicity with accuracy of $\\Delta M \\leq $ 0.3 mag. Our target in near future would be to adopt this procedure to RC stars."
    },
    "1208/1208.0489_arXiv.txt": {
        "abstract": "The cosmic star formation rate density (CSFRD) has been observationally investigated out to redshift $z\\simeq 10$.  However, most of the theoretical models for galaxy formation underpredict the CSFRD at $z\\gtrsim 1$.  Since the theoretical models reproduce the observed luminosity functions (LFs), luminosity densities (LDs), and stellar mass density at each redshift, this inconsistency does not simply imply that theoretical models should incorporate some missing unknown physical processes in galaxy formation.  Here, we examine the cause of this inconsistency at UV wavelengths by using a mock catalog of galaxies generated by a semi-analytic model of galaxy formation. We find that this inconsistency is due to two observational uncertainties: the dust obscuration correction and the conversion from UV luminosity to star formation rate (SFR).  The methods for correction of obscuration and SFR conversion used in observational studies result in the overestimation of the CSFRD by $\\sim 0.1$--$0.3$~dex and $\\sim 0.1$--$0.2$~dex, respectively, compared to the results obtained directly from our mock catalog.  We present new empirical calibrations for dust attenuation and conversion from observed UV LFs and LDs into the CSFRD. ",
        "introduction": "\\label{section:intro}  The cosmic star formation rate density (CSFRD), $\\dot{\\rho}_\\star$, is one of the most fundamental quantities that reveals how present galaxies were formed and evolved in the universe.  It has been probed observationally since the seminal works of Lilly et al. (1996) and Madau et al. (1996, 1998) using various star formation rate (SFR) indicators such as luminosities of the stellar continuum at the rest-frame ultraviolet (UV) and nebular emission lines (e.g., H$\\alpha$).  Various observational results are compiled in Hopkins (2004; H04) and Hopkins \\& Beacom (2006; HB06), in which the cosmology, stellar initial mass function (IMF), and dust obscuration correction are unified.  The best-fit CSFRD function using these results has been widely utilized not only in observational studies (e.g., Karim et al. 2011) but also in theoretical studies (e.g., Coward et al. 2008; Kistler et al. 2009; Tominaga et al. 2011; Wang \\& Dai 2011).  The CSFRD at $z\\lesssim 1$ has been confirmed by various SFR indicators obtained from wide-field surveys such as the \\textit{Galaxy Evolution Explorer} (e.g., Wyder et al. 2005; Robotham \\& Driver 2011).  Its characteristic feature is a rapid increase with redshift ($\\dot{\\rho}_\\star \\propto (1+z)^{3-4}$ at $z\\lesssim 1$; H04 and references therein).  Thus, the CSFRD at $z\\sim 1$ is an order of magnitude higher than that in the local universe.  In contrast, the CSFRD is still uncertain in the higher redshift range (i.e., $z\\gtrsim 1$) where popular SFR indicators include the rest-frame UV continuum stellar emission and infrared (IR) dust emission.  This uncertainty is due to uncertainties in the estimation of the CSFRD from observed data: the dust obscuration correction for the UV continuum, contamination from the old stellar population to the IR luminosity, estimation of the total IR luminosity, the faint-end slope of the luminosity function (LF), and the conversion factor from luminosity into SFR.  These uncertainties result in the well-known inconsistencies in some physical quantities between direct measurements and values inferred from the HB06 CSFRD function such as stellar mass density (SMD; $\\rho_\\star$: e.g., Wilkins et al. 2008; Choi \\& Nagamine 2012; Benson 2012), core-collapse supernova rate (e.g., Horiuchi et al. 2011; see also Botticella et al. 2012 who report there is no inconsistency in the local 11 Mpc volume), and extragalactic background light (e.g., Raue \\& Meyer 2012).  The CSFRD has also been calculated theoretically using galaxy formation models: hydrodynamic simulation (e.g., Nagamine et al. 2006) or semi-analytic models (e.g., Cole et al. 2001; Nagashima \\& Yoshii 2004; Benson 2012).  In these models, a galaxy-by-galaxy basis calculation is executed based on a detailed hierarchical structure formation scenario.  Therefore, the CSFRD at a certain redshift can be calculated by simply integrating the SFR of each galaxy at that redshift.  Although these theoretical models reproduce reasonably well the observed LFs and the luminosity densities (LDs) at the rest-frame wavelength dominated by stellar emission, and the SMD at both local and high $z$, most underpredict the CSFRD compared to that estimated observationally (e.g., Nagashima \\& Yoshii 2004; Nagashima et al. 2005; Lacey et al. 2011; Benson 2012).  While the underprediction of the CSFRD might be attributed to some missing unknown physical processes of galaxy formation and evolution, it is worth investigating whether or not the uncertainties in estimating the CSFRD from directly observed data could be the origin of the disagreement.  \\begin{figure*} \\epsscale{0.73} \\plotone{f1.eps} \\caption{ Redshift evolution of the SMD (top), the rest-frame UV (i.e., $\\lambda = 1500$--2800~{\\AA}) LD (middle), and the CSFRD (bottom). The solid curves in each panel are the predictions of the Mitaka model.  The thin solid curve in the top panel is the SMD for model galaxies with $M_\\star \\ge 10^8\\ M_\\odot$.  The dashed curve in the bottom panel is the best-fit Cole et al. (2001) functional form to the H04 $\\dot{\\rho}_\\star$.  The symbols with error bars are observational data.  The observational data shown in the top panel are compiled in Yabe et al. (2009) and those in the middle and bottom panels are from H04, respectively.  In the bottom panel, the boxes (circles) are evaluated using a common (SFR-dependent) obscuration correction from H04 for the same observed data plotted in the middle panel. } \\label{fig-LD+CSFR} \\end{figure*} Here, we examine the CSFRD through a comparison of the observational data compiled in H04 with the mock catalog of galaxies generated by one of the semi-analytic models for galaxy formation, the so-called Mitaka model (Nagashima \\& Yoshii 2004; see also Nagashima et al. 2005).  The Mitaka model reproduces various kinds of observations not only for local galaxies including the stellar continuum LFs (Nagashima \\& Yoshii 2004) but also for high-$z$ Lyman-break galaxies and Ly$\\alpha$ emitters (Kashikawa et al. 2006; Kobayashi et al. 2007, 2010).  In this paper, we focus on the rest-frame UV continuum luminosity as an SFR indicator; this is usually applied to galaxies in the high-$z$ universe (e.g., Madau et al. 1996, 1998). Investigation of other SFR indicators, such as the rest-frame IR continuum emitted by interstellar dust, will be the subject of future work.  As shown in the top and middle panels of Figure~\\ref{fig-LD+CSFR}, the Mitaka model reproduces well the measured SMD, $\\rho_\\star$, and LD at rest-frame UV wavelengths, $\\rho_\\mathrm{UV}$, which have not been corrected for interstellar dust attenuation, in the redshift range $z = 0$--$6$.  However, the model prediction for the CSFRD is underestimated at $z\\gtrsim 1$ by a factor of $\\sim 2$--3 (i.e., $\\sim 0.3$--0.5~dex) relative to the median value of $\\dot{\\rho}_\\star$ compiled in H04 as shown in the bottom panel of Figure~\\ref{fig-LD+CSFR}.  It should be emphasized that these H04 CSFRD data are calculated using the same observational UV LDs plotted in the middle panel of Figure~\\ref{fig-LD+CSFR}, and which are reasonably reproduced by our model.  Moreover, our CSFRD is consistent with the upper limit for $\\dot{\\rho}_\\star$ given by Strigari et al. (2005) estimated from an upper limit on the diffuse supernova neutrino background with Super-Kamiokande.  In this paper we show that the underestimation of the Mitaka model relative to the H04 CSFRD can be fully attributed to two observational uncertainties, the dust obscuration correction and the SFR conversion used in observational studies.  Therefore the underestimation using theoretical galaxy formation models relative to the H04 CSFRD does not necessarily imply that these models should incorporate some missing unknown physical processes.  \\begin{figure*} \\epsscale{.98} \\plotone{f2.eps} \\caption{ LD contribution from galaxies brighter than the horizontal axis of $L/L_\\ast$ calculated using the Schechter function with various faint-end slopes of the LF.  As on the figure, the solid, long-dashed, short-dashed, dotted, and dash-dotted curves are for $\\alpha = -1.2$, $-1.4$, $-1.6$, $-1.8$, and $-1.9$, respectively. The left and right panels are the LD normalized by $\\phi_\\ast L_\\ast$ and the total LD, respectively. } \\label{fig-Schechter} \\end{figure*} The paper is organized as follows.  In Section~\\ref{sec:ObsUnc}, we list the uncertainties in estimating the CSFRD from observed UV LFs. In Section~\\ref{sec:Model}, we describe the key prescriptions of the Mitaka model relating to the observational uncertainties in estimating the CSFRD.  Then we compare the model results with the observed LFs, LDs, and CSFRD in the redshift range $z=0$--10 in Section~\\ref{sec:Results}.  We summarize our work in Section~\\ref{sec:Summary+Discussion}, where we also provide a discussion which includes a new formula for the obscuration correction at rest-frame UV wavelengths in Section~\\ref{sec:Summary+Discussion}.  Throughout this paper, we adopt the 737 cosmology, i.e., $H_0=70~\\mathrm{km\\ s^{-1}\\ Mpc^{-1}}$ (i.e., $h_{70}\\equiv h/0.7 = 1$), $\\Omega_M=0.3$, and $\\Omega_\\Lambda =0.7$, and a Salpeter IMF with a mass range of $0.1$--$60~M_\\odot$. All magnitudes are expressed in the AB system and the wavelengths are given in the rest frame unless otherwise stated. ",
        "conclusions": "\\label{sec:Summary+Discussion}  In this paper, we examined the cause of the inconsistency in the CSFRD between theoretical and observational studies, focusing on an SFR indicator for a high-$z$ universe, that is, the rest-frame 1500~{\\AA} stellar continuum luminosity $L_{1500}$.  By using a semi-analytic model for galaxy formation, the so-called Mitaka model, we found that the underestimation of the CSFRD seen in theoretical model originates from the following two uncertainties in the process to evaluate the CSFRD from the observed 1500~{\\AA} LD: the dust obscuration correction and the conversion from $L_{1500}$ to SFR. The uncertainty in the faint-end slope of the LF is not the origin of the underestimation of CSFRD but the origin of the dispersion around the median CSFRD.  The methods of obscuration correction adopted in H04 result in the overestimation of the CSFRD by $\\approx 0.1$--0.4~dex and the SFR conversion used in observational studies also leads to the overestimation of the CSFRD by $\\approx 0.1$--0.2~dex.  Since theoretical models including ours reproduce the observed data for the UV LD which is not corrected by dust attenuation and the data for the SMD, the inconsistency in the CSFRD does not imply that theoretical models miss some key physical processes in galaxy formation.  Of course, the theoretical models are not yet perfect because observed data such as cosmic downsizing (e.g., Cowie et al. 1996) remain to be reproduced.  The revision of theoretical models can be achieved through a comparison with direct observed data which are \\textit{not} affected by a certain model and/or assumptions.  In this section, we provide a brief discussion of the origin of the difference in dust attenuation at high $z$ between our model and the H04 corrections.  We also present new empirical calibrations for dust attenuation and SFR conversion as well as a recipe for utilizing them in observational studies.  \\subsection{Origin of the Difference in Dust Attenuation at High Redshift}  As described in Section~\\ref{subsec:result-dust}, the H04 obscuration correction methods reproduce the intrinsic LD from the observable UV LF of our model galaxies at $z=0$, although they overestimate it at higher redshifts.  Here we discuss the origin of the overestimation at high redshift.  Our model naturally incorporates the redshift evolution of dust abundance because it calculates chemical enrichment of gas in each model galaxy consistently according to its SFH.  This can be seen in Figure~\\ref{fig-Alameff-z}, which shows the redshift evolution of \\textit{effective} dust attenuation in magnitude, $A_{1500}^\\mathrm{eff}$, for our model galaxies defined by \\begin{equation} A_{1500}^\\mathrm{eff}(z) \\equiv -2.5 \\log_{10}{\\left(\\rho_{1500}(z) / \\rho_{1500}^\\mathrm{int}(z) \\right)}. \\label{eq-Alameff} \\end{equation} \\begin{figure} \\epsscale{1.15} \\plotone{f10.eps} \\caption{ Same as Figure~\\ref{fig-CSFReff-z}, but for the ratio of the observable 1500~{\\AA} LD to the intrinsic one (i.e., $A_{1500}^\\mathrm{eff}$) of the Mitaka model as a function of redshift.  The parameters for the best-fit analytic function defined by Equation~(\\ref{eq-CdustEff}) and represented as the thick solid curve are also given in Table~\\ref{tab-FitParameters}. } \\label{fig-Alameff-z} \\end{figure}  As $A_{1500}^\\mathrm{eff}$ can be rewritten using the $A_{1500}$ and $L_{\\nu,1500}^\\mathrm{int}$ for each galaxy as $A_{1500}^\\mathrm{eff} = -2.5\\log_{10} ( \\sum ( 10^{-0.4A_{1500}^i} \\times L_{\\nu, 1500}^{\\mathrm{int}, i} ) / $ $\\sum L_{\\nu,1500}^{\\mathrm{int},i})$, $A_{1500}^\\mathrm{eff}$ represents the mean dust attenuation weighted by the intrinsic UV continuum luminosity $L_{\\nu, 1500}^\\mathrm{int}$.  This is the reason why $A_{1500}^\\mathrm{eff}$ at $z=0$ ($= 1.3$~mag) is larger than $\\langle A_{1500} \\rangle$ at $z=0$ ($\\lesssim 1$~mag). $A_{1500}^\\mathrm{eff}$ becomes smaller at higher $z$ because of the redshift evolution of metallicity and dust abundance.  However, such a redshift evolution of dust abundance is not incorporated into the H04 obscuration corrections.  This is the origin of their overestimation of the intrinsic LD.  A quantity similar to $A_{1500}^\\mathrm{eff}$ has been evaluated from the observed LD ratio between IR and UV, $\\rho_\\mathrm{IR} / \\rho_\\mathrm{UV}$ ($\\approx \\rho_{1500}^\\mathrm{int} / \\rho_{1500} - 1$), in the redshift range $z=0$--1 (Takeuchi et al. 2005).  They found that $A_{1500}^\\mathrm{eff}$ increases monotonically toward $z=1$, which is the opposite trend to our result.  However, more recent observational estimates for $\\rho_\\mathrm{IR}$ (e.g., Murphy et al. 2011; Casey et al. 2012; Cucciati et al. 2012) find that the redshift evolution of $\\rho_\\mathrm{IR}$ in this redshift range is milder than that reported by Takeuchi et al. (2005).  Cucciati et al. (2012) also find that mean dust attenuation decreases toward high $z$ at $z\\gtrsim 1$; their result is consistent with our prediction.  These results may indicate that our model does not underestimate dust attenuation of galaxies at $z\\gtrsim 1$ but overestimates it at $z\\approx 0$.  This interpretation will be examined as in our future work.   \\subsection{Comparison with the HB06 CSFRD} \\label{subsec:Comments2HB06}  As described in Section~\\ref{subsec:CorrDustAtten}, HB06 obtained the CSFRD at $z < 3$ by summing the UV data and FIR measurements. They reported that this technique of UV$+$FIR measurements gives an effective obscuration correction to the UV data by a factor of 2 at $z\\approx 0$ and $\\sim 5$ (i.e., 1.7 mag) at $z \\gtrsim 1$.  It may be difficult to reconcile this with our result and should be examined in detail.  However, the correction factor of $\\sim 5$ at $z\\gtrsim 1$ reported in HB06 might be overestimated for the following reason.  HB06 performed the effective obscuration correction by just adding a constant CSFRD measured by Le Floc'h et al. (2005) for FIR wavelengths at $z=1$.  This is based on the observational result that FIR measurements of the CSFRD are quite flat in the range $z=1$--3 as reported by P\\'{e}rez-Gonz\\'{a}lez et al. (2005), who estimated the total IR luminosity of each galaxy by using the local template spectral energy distribution of Chary \\& Elbaz (2001). However, as reported by Murphy et al. (2011), using this template results in an overestimation of the total IR luminosity for IR-bright galaxies at $z>1.5$.  Hence, the IR measurement of P\\'{e}rez-Gonz\\'{a}lez et al. (2005) may be overestimated.  We are planning to investigate whether our model reproduces the observed IR data in future studies.   \\subsection{Empirical Calibrations of the Obscuration Correction and the SFR Conversion from UV Luminosity} \\label{subsec:empirical-calib}  Here we propose new empirical formulas which correct for dust obscuration and convert from the intrinsic UV LD, $\\rho_\\mathrm{UV}^\\mathrm{int}$, to the CSFRD, $\\dot{\\rho}_\\star$. These formulas are derived to reproduce the true quantities for the model galaxies and are represented by explicit analytic functions. For the obscuration correction, we show two different formulas; one is for the observable LF and the other is for its integrated quantity $\\rho_\\mathrm{UV}$.  These two formulas are, respectively, similar to the SFR-dependent and common corrections of H04, but they are described to have redshift dependence.  \\subsubsection{Conversion from Observable LF into Intrinsic LF}  Let us define $C_\\mathrm{dust}$ as an empirical formula to convert the observable LF into an intrinsic one.  $C_\\mathrm{dust}$ at a certain magnitude $M_{1500}$ is determined via an abundance-matching approach.  That is, the cumulative number density of the observable LF of the Mitaka model, $n^\\mathrm{obs}(< M_{1500})$, should match that of the intrinsic LF, $n^\\mathrm{int}(< M_{1500} - C_\\mathrm{dust})$: \\begin{equation} n^\\mathrm{obs}(< M_{1500}) = n^\\mathrm{int}(< M_{1500} - C_\\mathrm{dust}). \\end{equation} The top panel of Figure~\\ref{fig-Cdust-z0} shows the resultant $C_\\mathrm{dust}$ as a function of $M_{1500}$ at $z=0$. \\begin{figure} \\epsscale{1.15} \\plotone{f11.eps} \\caption{ Top, middle, and bottom panels show $C_\\mathrm{dust}$, LF, and the ratio of the LD calculated from the Mitaka observable LF $+C_\\mathrm{dust}$ to the Mitaka intrinsic LD, respectively, at $z=0$ for a wavelength of $\\lambda = 1500$~{\\AA}.  In the top panel, the filled circles are the numerical data calculated from the intrinsic and observable 1500~{\\AA} LFs of the Mitaka model, while the dashed curve is analytical fit using Equation~(\\ref{eq-Cdust-fit}) with the numerical quantities in Table~\\ref{tab-Cdust}.  $C_\\mathrm{dust} = 0$ at $M_{1500} \\lesssim -22$~mag is reflected by the fact that there is no model galaxy with such bright observable magnitudes.  In the middle panel, the solid and dotted curves are the intrinsic and observable LFs of the Mitaka model, respectively, while the dashed curve is the Mitaka observable LF + $C_\\mathrm{dust}$. } \\label{fig-Cdust-z0} \\end{figure}  We have derived $C_\\mathrm{dust}$ for other redshifts and found that $C_\\mathrm{dust}$ can be fitted well with the following analytic function in the redshift range $z=0$--10: \\begin{equation} C_\\mathrm{dust}(M_{1500};\\ z) = a\\ \\exp{\\left[-b |M_{1500} - M_{1500}^0|^c\\right]}. \\label{eq-Cdust-fit} \\end{equation} Here $a,\\ b,\\ c,$ and $M_{1500}^0$ are model parameters and evolve with redshift.  $a$ and $M_{1500}^0$ have units of magnitude, while $b$ and $c$ are non-dimensional constants.  We adopt a smoothly declining functional form even for the bright magnitudes where there are no model galaxies.  We note that $C_\\mathrm{dust}$ is a dust obscuration correction for the UV LF as a statistical quantity like the SFR-dependent correction of H04.  Hence it does not represent a mean dust attenuation for the galaxies at a certain magnitude.  Actually, the intrinsically UV-brightest galaxies are not the observationally brightest galaxies in our model, as described in Section 4.1.  The reason why $C_\\mathrm{dust}$ has a peak around the characteristic magnitude $M_{1500}^\\ast$ is that the magnitude difference between the intrinsic and observable UV LFs becomes largest at $M_{1500} \\approx M_{1500}^\\ast$, not that the mean attenuation of the galaxies at $M_{1500} \\approx M_{1500}^\\ast$ is the largest.  There are significant discrepancies between the raw $C_\\mathrm{dust}$ quantities and that evaluated using Equation~(\\ref{eq-Cdust-fit}), as represented by the filled circles and dashed curve, respectively, in the top panel of Figure~\\ref{fig-Cdust-z0} at $M_{1500} \\lesssim -23$~mag or $M_{1500} \\gtrsim -15$~mag.  Fortunately, this deviation hardly affects the corrected UV LF or the integrated LD as shown in the middle and bottom panels of Figure~\\ref{fig-Cdust-z0}. Table~\\ref{tab-Cdust} gives the numerical quantities for the best-fit parameters in the redshift range $z = 0$--10. \\begin{deluxetable}{ccccc} \\tablewidth{250pt} \\tabletypesize{\\scriptsize} \\tablecolumns{5} \\tablecaption{Fitting Parameters for our Formula to Correct Dust Attenuation, $C_\\mathrm{dust}$} \\tablehead{ Redshift, $z$ & $a$ & $b$ & $c$ & $M_{1500}^0$\\\\ & (mag) & & & (mag) } \\startdata 0  & $1.847$ & $0.1139$ & $1.566$ & $-19.49$ \\\\ 1  & $2.466$ & $0.1181$ & $1.281$ & $-22.65$ \\\\ 2  & $2.344$ & $0.2625$ & $1.086$ & $-21.86$ \\\\ 3  & $2.235$ & $0.2844$ & $1.184$ & $-21.91$ \\\\ 4  & $2.032$ & $0.2115$ & $1.428$ & $-22.15$ \\\\ 5  & $1.822$ & $0.1971$ & $1.599$ & $-22.16$ \\\\ 6  & $1.583$ & $0.1673$ & $1.737$ & $-22.14$ \\\\ 7  & $1.496$ & $0.1630$ & $1.800$ & $-22.08$ \\\\ 8  & $1.415$ & $0.1900$ & $1.641$ & $-22.16$ \\\\ 9  & $1.228$ & $0.1239$ & $1.895$ & $-22.39$ \\\\ 10 & $1.109$ & $0.1898$ & $1.634$ & $-22.17$ \\enddata \\tablecomments{The analytic expression for $C_\\mathrm{dust}$ is given by Equation~(\\ref{eq-Cdust-fit}).} \\label{tab-Cdust} \\end{deluxetable}  \\begin{deluxetable}{cccccc} \\tabletypesize{\\scriptsize} \\tablewidth{\\linewidth} \\tablecolumns{6} \\tablecaption{Fitting Parameters for our Formulas for $C_\\mathrm{dust}^\\mathrm{eff}$ and $C_\\mathrm{SFR}^\\mathrm{eff}$\\label{tab-FitParameters}} \\tablehead{ \\multicolumn{2}{c}{\\scriptsize $C_\\mathrm{dust}^\\mathrm{eff}$} & \\colhead{} & \\multicolumn{3}{c}{\\scriptsize $C_\\mathrm{SFR}^\\mathrm{eff}$}\\\\ \\cline{1-2} \\cline{4-6}\\\\ {\\scriptsize $\\alpha$} & {\\scriptsize $\\beta$} & \\colhead{} & {\\scriptsize $C_2$} & {\\scriptsize $C_1$} & {\\scriptsize $C_0$} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{} & {\\scriptsize ($M_\\odot \\ \\mathrm{yr^{-1}\\ (ergs/ s/ Hz)^{-1}}$)} } \\startdata {\\scriptsize $2.983$} & {\\scriptsize $0.3056$} & & {\\scriptsize $-5.915\\times 10^{-5}$} & {\\scriptsize $7.294\\times 10^{-4}$} & {\\scriptsize $-28.01$} \\enddata \\tablecomments{The analytic expressions for $C_\\mathrm{dust}^\\mathrm{eff}$ and $C_\\mathrm{SFR}^\\mathrm{eff}$ are given by Equations~(\\ref{eq-CdustEff}) and (\\ref{eq-Lnu2SFR}), respectively.} \\end{deluxetable}  The normalization factor $a$ indicates the maximum value of $C_\\mathrm{dust}$ at the redshift.  $a$ is found to gradually increase with redshift toward its peak at $z\\simeq 1$--3 and then decreases as $z$ increases.  It is interesting that the peak redshift for $a$ is roughly equal to the redshift where dusty galaxies (e.g., ultraluminous infrared galaxies, sub-mm galaxies, etc.) are mainly found.  \\subsubsection{Conversion from Observable UV LD into Intrinsic UV LD}  The conversion factor from the observed 1500~{\\AA} LD, $\\rho_{1500}$, into the intrinsic one, $\\rho_{1500}^\\mathrm{int}$, is defined as $C_\\mathrm{dust}^\\mathrm{eff} \\equiv \\rho_{1500}^\\mathrm{int} / \\rho_{1500}$.  This relates to the effective dust attenuation $A_{1500}^\\mathrm{eff}$ defined in Equation~(\\ref{eq-Alameff}) via $C_\\mathrm{dust}^\\mathrm{eff} = \\mathrm{dex}( 0.4 A_{1500}^\\mathrm{eff} )$.\\footnote{$\\mathrm{dex}(x)$ is the inverse function of $\\log_{10}(x)$: $\\mathrm{dex}(x) \\equiv 10^x$.}  It is found that the $C_\\mathrm{dust}^\\mathrm{eff}$ for our model galaxies can be well fitted by the following simple analytic function with two redshift-independent parameters $\\alpha$ and $\\beta$ in the redshift range $z = 0$--10 as shown in Figure~\\ref{fig-Alameff-z}: \\begin{equation} C_\\mathrm{dust}^\\mathrm{eff}(z) = \\alpha \\ \\exp{ \\left[-\\beta \\left(1+z \\right)\\right]} + 1. \\label{eq-CdustEff} \\end{equation} This functional form is motivated by the natural expectation that $C_\\mathrm{dust}^\\mathrm{eff}$ approaches unity at high redshift. Since there is little dust at high redshift, $C_\\mathrm{dust}^\\mathrm{eff} \\approx 1$.  The best-fit parameters which reproduce the ratio $\\rho_{1500}^\\mathrm{int} / \\rho_{1500}$ at $z=0$--10 within $\\pm 10\\%$ are given in Table~\\ref{tab-FitParameters}.  While the normalization parameter $a$ in Equation~(\\ref{eq-Cdust-fit}) has its peak at $z\\simeq 1$--3, $C_\\mathrm{dust}^\\mathrm{eff}$ decreases monotonically with increasing $z$.  The different redshift evolution can be interpreted by the fact that the contribution from galaxies fainter than $M_{1500}^0$, where the galaxies have a maximum extinction of $a$ at the redshift, to the intrinsic 1500~{\\AA} LD is significant; for a typical faint-end slope $\\alpha \\approx -1.4$ of the UV LF, the contribution from galaxies with $L \\lesssim L_\\ast $ reaches $\\sim 0.7$~dex as shown in the right panel of Figure~\\ref{fig-Schechter}.  Since such faint galaxies have smaller $C_\\mathrm{dust}$ than its peak value of $a$, $C_\\mathrm{dust}^\\mathrm{eff}$ progressively decreases toward high redshift although $a$ has its peak at $z\\simeq 1$--3.  \\subsubsection{Conversion from Intrinsic UV LD into CSFRD}  We find that the ratio of the CSFRD to the intrinsic 1500~{\\AA} LD, $\\dot{\\rho}_\\star / \\rho_{1500}^\\mathrm{int}$, can be well fitted with the following simple quadratic function in the redshift range $z=0$--10: \\begin{equation} C_\\mathrm{SFR}^\\mathrm{eff}(z) = C_0\\left[1 + C_1\\left(1+z\\right) + C_2\\left(1+z\\right)^2\\right]. \\label{eq-Lnu2SFR} \\end{equation} Here $C_0,\\ C_1,$ and $C_2$ are the model parameters and are redshift-independent constants.  $C_0$ has the same dimensions as $C_\\mathrm{SFR}^\\mathrm{eff}$, $M_\\odot\\ \\mathrm{yr^{-1}\\ (erg\\ s^{-1}\\ Hz^{-1})^{-1}}$, while $C_1$ and $C_2$ are non-dimensional constants.  With the best-fit quantities given in Table~\\ref{tab-FitParameters}, the analytic function reproduces our model results for $\\dot{\\rho}_\\star / \\rho_{1500}^\\mathrm{int}$ within $\\pm 5$\\%.  We note here that, while the empirical formula given in Equation~(\\ref{eq-Lnu2SFR}) predicts that $C_\\mathrm{SFR}^\\mathrm{eff}$ has a peak at $z\\sim 6$ and progressively decreases toward high redshifts, it is simply a fitting result and does not have any physical motivation. Nevertheless, it might be a real trend as discussed in Section~\\ref{subsec-ResCSFR}.  \\subsection{Recipe for Converting Observed UV LF into CSFRD}  Here we describe a recipe for converting the observed 1500~{\\AA} LF into the CSFRD $\\dot{\\rho}_\\star$ using our empirical formulas given in Section~\\ref{subsec:empirical-calib}.  As a first step, the intrinsic 1500~{\\AA} LD, $\\rho_{1500}^\\mathrm{int}$, should be calculated from the observed UV LF.  This can be done using either one of the following two approaches.  The first approach is to convert the observed UV LF into an intrinsic LF via the empirical formula for $C_\\mathrm{dust}$ as a function of magnitude and redshift given in Equation~(\\ref{eq-Cdust-fit}) with the best-fit parameters in Table~\\ref{tab-Cdust}.  Then the intrinsic UV LF is integrated over magnitude to obtain the intrinsic UV LD, $\\rho_{1500}^\\mathrm{int}$.  The second approach is to first integrate the observed UV LF over magnitude to obtain the observable UV LD, $\\rho_{1500}$, and then converted it into $\\rho_{1500}^\\mathrm{int}$ using the empirical formula for $C_\\mathrm{dust}^\\mathrm{eff}$ as a function of redshift given in Equation~(\\ref{eq-CdustEff}).  Finally, one can evaluate $\\dot{\\rho}_\\star$ from $\\rho_{1500}^\\mathrm{int}$ using the empirical formula for $C_\\mathrm{SFR}^\\mathrm{eff}$ as a function of redshift given in Equation~(\\ref{eq-Lnu2SFR}).  For $C_\\mathrm{dust}$ given in Equation~(\\ref{eq-Cdust-fit}), it is statistically enough to linearly interpolate the formula for a specified redshift while the best-fit parameters are provided discretely in redshift for our empirical formula.  One should interpolate the parameters to evaluate an adequate $C_\\mathrm{dust}$ at a certain magnitude and the desired redshift rather than interpolating $C_\\mathrm{dust}$ itself.  \\begin{deluxetable*}{ccccc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecolumns{5} \\tablecaption{LD and CSFRD of the Mitaka Model} \\tablehead{ Redshift, $z$ & $\\dot{\\rho_\\star}$ & $\\rho_{1500}$ & $\\rho_{1500}^\\mathrm{int}$ & $\\dot{\\rho_\\star} / \\rho_{1500}$\\\\ & $(h_{70}\\ M_\\odot\\ \\mathrm{yr^{-1}\\ Mpc^{-3}})$ & $(h_{70}\\ \\mathrm{erg\\ s^{-1}\\ Hz^{-1}\\ Mpc^{-3}})$ & $(h_{70}\\ \\mathrm{erg\\ s^{-1}\\ Hz^{-1}\\ Mpc^{-3}})$ & $(M_\\odot\\ \\mathrm{yr^{-1}\\ (erg\\ s^{-1}\\ Hz^{-1})^{-1}})$} \\startdata 0 & $2.599\\times 10^{-2}$ & $8.090\\times 10^{25}$ & $2.616\\times 10^{26}$ & $3.213\\times 10^{-28}$ \\\\ 1 & $6.825\\times 10^{-2}$ & $2.531\\times 10^{26}$ & $6.593\\times 10^{26}$ & $2.697\\times 10^{-28}$ \\\\ 2 & $8.285\\times 10^{-2}$ & $3.665\\times 10^{26}$ & $7.910\\times 10^{26}$ & $2.261\\times 10^{-28}$ \\\\ 3 & $6.828\\times 10^{-2}$ & $3.498\\times 10^{26}$ & $6.396\\times 10^{26}$ & $1.952\\times 10^{-28}$ \\\\ 4 & $4.926\\times 10^{-2}$ & $2.752\\times 10^{26}$ & $4.518\\times 10^{26}$ & $1.790\\times 10^{-28}$ \\\\ 5 & $3.124\\times 10^{-2}$ & $1.902\\times 10^{26}$ & $2.818\\times 10^{26}$ & $1.642\\times 10^{-28}$ \\\\ 6 & $1.759\\times 10^{-2}$ & $1.119\\times 10^{26}$ & $1.540\\times 10^{26}$ & $1.572\\times 10^{-28}$ \\\\ 7 & $9.105\\times 10^{-3}$ & $5.988\\times 10^{25}$ & $7.849\\times 10^{25}$ & $1.521\\times 10^{-28}$ \\\\ 8 & $4.524\\times 10^{-3}$ & $3.129\\times 10^{25}$ & $3.865\\times 10^{25}$ & $1.446\\times 10^{-28}$ \\\\ 9 & $2.069\\times 10^{-3}$ & $1.641\\times 10^{25}$ & $1.898\\times 10^{25}$ & $1.261\\times 10^{-28}$ \\\\ 10& $9.514\\times 10^{-4}$ & $8.245\\times 10^{24}$ & $9.150\\times 10^{24}$ & $1.154\\times 10^{-28}$ \\enddata \\tablecomments{$\\dot{\\rho}_\\star$ is the CSFRD and $\\rho_{1500}$ and $\\rho_{1500}^\\mathrm{int}$ represent the observable and intrinsic 1500~{\\AA} LDs, respectively, for all of the galaxies in the Mitaka model.} \\label{tab-CSFRDdata} \\end{deluxetable*}  \\begin{deluxetable*}{lccr} \\tabletypesize{\\scriptsize} \\tablecolumns{4} \\tablecaption{CSFRD Parametric Fits to Various Forms from the Literatures} \\tablehead{ Reference & Functional Form & Parameter & Value } \\startdata Cole et al. (2001) & $\\dot{\\rho}_\\ast(z) = \\left(a+bz\\right)h\\ /\\ [1+\\left(z/c\\right)^d]$ & $a$ & $0.0389$ \\\\ & & $b$ & $0.0545$ \\\\ & & $c$ & $2.973$ \\\\ & & $d$ & $3.655$ \\\\ \\hline Hernquist \\& Springel (2003) & $\\dot{\\rho}_\\ast(z) = \\dot{\\rho}_0\\chi^2\\ /\\ [1+\\alpha (\\chi-1)^3\\exp{(\\beta \\chi^{7/4})}]$ & $\\dot{\\rho}_0$ & $0.030$\\\\ & $\\chi=[H(z)/H_0]^{2/3}$ & $\\alpha$ & $0.323$\\\\ & & $\\beta$ & $0.051$\\\\ \\hline Y\\\"{u}ksel et al. (2008) & $\\dot{\\rho}_\\ast(z) = \\dot{\\rho}_0 [(1+z)^{\\alpha \\eta}+\\{(1+z)/B\\}^{\\beta \\eta}+\\{(1+z)/C\\}^{\\gamma \\eta}]^{1/\\eta}$, & $\\dot{\\rho}_0$ &  $0.0258$ \\\\ & $B=(1+z_1)^{1-\\alpha /\\beta}$, \\phm{...}$C=(1+z_1)^{(\\beta -\\alpha)/\\gamma}(1+z_2)^{1-\\beta/\\gamma}$ & $\\alpha$   &  $1.6$ \\\\ & & $\\beta$  & $-1.2$ \\\\ & & $\\gamma$ & $-5.7$ \\\\ & & $z_1$ &  $1.7$ \\\\ & & $z_2$ &  $5.0$ \\\\ & & $\\eta$ & $-1.62$ \\enddata \\label{tab-CSFRDfits} \\end{deluxetable*}  The numerical quantities for $\\dot{\\rho}_\\star,\\ \\rho_{1500},$ and $\\rho_{1500}^\\mathrm{int}$ of the Mitaka model are compiled in Table~\\ref{tab-CSFRDdata}.  We also present the CSFRD parametric fits to a variety of analytic forms in the literature (Cole et al. 2001; Hernquist \\& Springel 2003; Y\\\"{u}ksel et al. 2008) in Table~\\ref{tab-CSFRDfits}."
    },
    "1208/1208.2997_arXiv.txt": {
        "abstract": "{What else can be said about star formation rate indicators that has not been said already many times over? The `coming of age' of large ground-based surveys and the unprecedented sensitivity, angular resolution and/or field-of-view of infrared and ultraviolet space missions have provided extensive, homogeneous data on both nearby and distant galaxies, which have been used to further our understanding of the strengths  and pitfalls of many common star formation rate indicators. The synergy between these surveys has also enabled the calibration of indicators for use on scales that are comparable to those of star-forming regions, thus much smaller than an entire galaxy. These are being used to investigate star formation processes at the sub-galactic scale. I review progress in the field over the past decade or so.} \\newline\\newline  \\def\\thechapter{1} ",
        "introduction": "My goal for this chapter, based on a series of lectures at the XXIII Canary Islands Winter School of Astrophysics, is to present  current understanding and calibrations of star formation rate (SFR) indicators, both on global, galaxy-wide scales, and on local, sub-galactic scales. SFRs are, together with masses, the most important parameters that define galaxies and their evolution across cosmic times. Although SFR calibrations have existed, with various levels of accuracy, for many years and sometimes decades, the past eight to ten years have brought forth major progress, through cohesive, multi-wavelength surveys of nearby and distant galaxies. These surveys have exploited the sensitivity, angular resolution and/or large field of view of  space telescopes (e.g., the \\textit{Spitzer Space Telescope}, the \\textit{Galaxy Evolution Explorer} [\\textit{GALEX}], the \\textit{Hubble Space Telescope} [\\textit{HST}], and the \\textit{Herschel Space Telescope}) and leveraged the multi-band coverage supplied by ground-based, all-sky surveys (e.g., the Sloan Digital Sky Survey), in order to push the definition of SFR indicators into new regimes, both in terms of wavelength coverage and spatial scales. I review here recent progress in this area, but also highlight where challenges, sometimes unexpected ones, have arisen. This chapter is structured to provide also a quick reference for the relevant literature on SFR calibrations. It loosely follows the structure of the lectures I presented at the Winter School. ",
        "conclusions": ""
    },
    "1208/1208.0676_arXiv.txt": {
        "abstract": "Observations of radio halos and relics in galaxy clusters indicate efficient electron acceleration. Protons should likewise be accelerated and, on account of weak energy losses, can accumulate, suggesting that clusters may also be sources of very high-energy (VHE; $E>100$ GeV) gamma-ray emission. We report here on VHE gamma-ray observations of the Coma galaxy cluster with the VERITAS array of imaging Cherenkov telescopes, with complementing \\Fermi-LAT observations at GeV energies. No significant gamma-ray emission from the Coma cluster was detected. Integral flux upper limits at the 99\\% confidence level were measured to be on the order of $(2-5)\\times 10^{-8}\\ {\\rm ph.\\,m^{-2}\\,s^{-1}}$ (VERITAS, $>220\\ {\\rm GeV}$) and $\\sim 2\\times 10^{-6}\\ {\\rm ph.\\,m^{-2}\\, s^{-1}}$ (\\Fermi, $1-3\\ {\\rm GeV}$), respectively. We use the gamma-ray upper limits to constrain CRs and magnetic fields in Coma. Using an analytical approach, the CR-to-thermal pressure ratio is constrained to be $< 16\\%$ from VERITAS data and $< 1.7\\%$ from \\Fermi data (averaged within the virial radius). {These upper limits are starting to constrain the CR physics in self-consistent cosmological cluster simulations and cap the maximum CR acceleration efficiency at structure formation shocks to be $<50\\%$. Alternatively, this may argue for non-negligible CR transport processes such as CR streaming and diffusion into the outer cluster regions. } Assuming that the radio-emitting electrons of the Coma halo result from hadronic CR interactions, the observations imply a lower limit on the central magnetic field in Coma of $\\sim (2 - 5.5)\\,\\mu{\\rm G}$, depending on the radial magnetic-field profile and on the gamma-ray spectral index. Since these values are below those inferred by Faraday rotation measurements in Coma (for most of the parameter space), this {renders} the hadronic model a very plausible explanation of the Coma radio halo. Finally, since galaxy clusters are dark-matter (DM) dominated, the VERITAS upper limits have been used to place constraints on the thermally-averaged product of the total self-annihilation cross section and the relative velocity of the DM particles, $\\expval{\\sigma v}$. ",
        "introduction": "Clusters of galaxies are the largest virialized objects in the Universe, with typical sizes of a few Mpc and masses on the order of $10^{14}$ to $10^{15} M_{\\odot}$. According to the currently favored hierarchical model of cosmic structure formation, larger objects formed through successive mergers of smaller objects with galaxy clusters sitting on top of this mass hierarchy \\citep[see][for a review]{article:Voit:2005}. Most of the mass ($\\sim$80\\%) in a cluster is dark matter (DM), as indicated by galaxy dynamics and gravitational lensing \\citep{article:DiaferioSchindlerDolag:2008}. Baryonic gas making up the intra-cluster medium (ICM) contributes about 15\\% of the total cluster mass and individual galaxies account for the remainder (about 5\\%). The ICM gas mass also comprises a significant fraction of the observable (baryonic) matter in the Universe.  The ICM is a hot ($T\\sim 10^{8}$ K) plasma emitting thermal bremsstrahlung in the soft X-ray regime \\citep[see, e.g.,][]{article:Petrosian:2001}. This plasma has been heated primarily through collisionless structure-formation shocks that form as a result of the hierarchical merging and accretion processes. Such shocks and turbulence in the ICM gas in combination with intra-cluster magnetic fields also provide a means to accelerate particles efficiently \\citep[see, e.g.,][]{article:ColafrancescoBlasi:1998, article:Ryu_etal:2003}. Many clusters feature megaparsec scale halos of nonthermal radio emission, indicative of a population of relativistic electrons and magnetic fields permeating the ICM \\citep{article:Cassano_etal:2010}. There are two competing theories to explain radio halos. In the ``hadronic model'', the radio-emitting electrons and positrons are produced in inelastic collisions of cosmic-ray (CR) ions with the thermal gas of the ICM \\citep{article:Dennison:1980, article:EnsslinPfrommerMiniatiSubramanian:2011}. In the ``re-acceleration model'', a long-lived pool of 100-MeV electrons---previously accelerated by formation shocks, galactic winds, or jets of active galactic nuclei (AGN)---interacts with plasma waves that are excited during states of strong ICM turbulence, e.g., after a cluster merger. This may result in second order Fermi acceleration and may produce energetic electrons ($\\sim 10$ GeV) sufficient to explain the observable radio emission \\citep{article:SchlickeiserSieversThiemann:1987, article:BrunettiLazarian:2010}. Observations of possibly nonthermal emission from clusters in the extreme ultraviolet \\citep[EUV; ][]{article:SarazinLieu:1998} and hard X-rays \\citep{article:RephaeliGruber:2002, article:Fusco-Femiano_etal:2004, article:Eckert_etal:2007} may provide further indication of relativistic particle populations in clusters, although the interpretation of these observations as nonthermal diffuse emission has been disputed on the basis of more sensitive observations \\citep[see, e.g.,][]{article:Ajello_etal:2009, article:Ajello_etal:2010, article:Wik_etal:2009}.  Galaxy clusters have, for many years, been proposed as sources of gamma rays. If shock acceleration in the ICM is an efficient process, a population of highly relativistic CR protons and heavy ions is to be expected in the ICM. The main energy-loss mechanism for CR hadrons at high energies is pion production through the interaction of CRs with nuclei in the ICM. Pions are short lived and decay. The decay of neutral pions produces gamma rays and the decay of charged pions produces muons, which then decay to electrons and positrons. Due to the low density of the ICM ($n_{\\mathrm{ICM}}\\sim 10^{-3}$ cm$^{-3}$), the large size and the volume-filling magnetic fields in the ICM, CR hadrons will be confined in the cluster on timescales comparable to, or longer than, the Hubble time \\citep[][]{article:Volk_etal:1996, article:Berezinsky_etal:1997} and they can therefore accumulate. For a given CR distribution function, the hadronically induced gamma-ray flux is directly proportional to the CR-to-thermal pressure fraction, $X_\\CR=\\expval{P_{\\CR}}/ \\expval{P_{\\mathrm{th}}}$ \\citep[see, e.g.,][]{article:EnsslinPfrommerSpringelJubelgas:2007}, where the brackets indicates volume averages. A very modest $X_{\\CR}$ of a few percent implies an observable flux of gamma rays \\citep[e.g.,][]{article:PfrommerEnsslin:2004b}.  Hydrostatic estimates of cluster masses, which are determined by balancing the thermal pressure force and the gravitational force, are biased low by the presence of any substantial nonthermal pressure component, including a CR pressure contribution. Similarly, a substantial CR pressure can bias the temperature decrement of the cosmic microwave background (CMB) due to the Sunyaev-Zel'dovich effect in the direction of a galaxy cluster. This could then severely jeopardize the use of clusters to determine cosmological parameters. Comparing X-ray and optical potential profiles in the centers of galaxy clusters yields an upper limit of 20-30\\% of nonthermal pressure (that can be composed of CRs, magnetic fields or turbulence) relative to the thermal gas pressure \\citep{article:Churazov_etal:2008, article:Churazov_etal:2010}. An analysis that compares spatially resolved weak gravitational lensing and hydrostatic X-ray masses for a sample of 18 galaxy clusters detects a deficit of the hydrostatic mass estimate compared to the lensing mass of $20\\%$ at $R_{500}$ -- the radius within which the mean density is 500 times the critical density of the Universe -- suggesting again a substantial nonthermal pressure contribution on large scales \\citep{article:Mahdavi_etal:2008}. Observing gamma-ray emission is a complementary method of constraining the pressure contribution of CRs that is most sensitive to the cluster core region. However, it assumes that the CR component is fully mixed with the ICM and may not allow for a detection of a two-phase structure of CRs and the thermal ICM. An $X_\\CR$ of only a few percent is required in order to produce a gamma-ray flux observable with the current generation of gamma-ray telescopes, rendering this technique at least as sensitive as the dynamical and hydrostatic methods (which are more general in that they are sensitive to any nonthermal pressure component).  Gamma-ray emission can also be produced by Compton up-scattering of ambient photons, for example CMB photons, on ultra-relativistic electrons. Those electrons can either be secondaries from the CR interactions mentioned above, or injected into the ICM by powerful cluster members and further accelerated by diffusive shock acceleration or turbulent reacceleration processes \\citep[][and references therein]{article:SchlickeiserSieversThiemann:1987}.  A third mechanism for gamma-ray production in a galaxy cluster could be self-annihilation of a DM particle, e.g., a weakly interacting massive particle (WIMP). As already mentioned, about 80\\% of the cluster mass is in the form of dark matter, which makes galaxy clusters interesting targets for DM searches \\citep{article:EvansFerrerSarkar:2004, article:BergstromHooper:2006, article:PinzkePfrommerBergstrom2009, article:Cuesta_etal:2011} despite their large distances compared to other common targets for DM searches, such as dwarf spheroidal galaxies \\citep{article:Strigari_etal:2007, article:Acciari_etal:2010, article:Aliu_etal:2009} or the Galactic Center \\citep{article:Kosack_etal:2004, article:Aharonian_etal:2006, article:Aharonian_etal:2009b, article:Abramowski_etal:2011}.  While several observations of clusters of galaxies have been made with satellite-borne and ground-based gamma-ray telescopes, a detection of gamma-ray emission from a cluster has yet to be made. Observations with EGRET \\citep{article:Sreekumar_etal:1996, article:Reimer_etal:2003} and the Large Area Telescope (LAT) on board the \\Fermi Gamma-ray Space Telescope \\citep{article:Ackermann_etal:2010} have provided upper limits on the gamma-ray fluxes (typically $\\sim10^{-9}$ ph cm$^{2}$ s$^{-1}$ for \\Fermi-LAT observations) for several galaxy clusters in the MeV to GeV band. Upper limits on the very-high-energy (VHE) gamma-ray flux from a small sample of clusters, including the Coma cluster, have been provided by observations with ground-based imaging atmospheric Cherenkov telescopes \\citep[IACTs;][]{article:Perkins_etal:2006, inproc:Perkins_etal:2008, article:Aharonian_etal:2009a, article:Aleksic_etal:2010,article:Aleksic_etal:2012}.  The Coma cluster of galaxies (ACO 1656) is one of the most thoroughly studied clusters across all wavelengths \\citep{article:Voges_etal:1999}. Located at a distance of about 100 Mpc \\citep[$z=0.023$;][]{article:StrubleRood:1999}, it is one of the closest massive clusters \\citep[$M \\sim10^{15}M_{\\odot}$;][]{article:Smith:1983, article:Kubo_etal:2008}. It hosts both a giant radio halo \\citep{article:Giovannini_etal:1993,article:Thierbach_etal:2003} and peripheral radio relic, which appears connected to the radio halo with a ``diffuse'' bridge \\citep[see discussion in][]{article:BrownRudnick:2010}. It has been suggested \\citep{article:Ensslin_etal:1998} and successively demonstrated by cosmological simulations which model the nonthermal emission processes \\citep{article:PfrommerEnsslinSpringel:2008, article:Pfrommer:2008, article:Battaglia_etal:2009, article:Skillman_etal:2011}, that the relic could well be an infall shock. Extended soft thermal X-ray (SXR) emission is evident from the ROSAT all-sky survey in the 0.1 to 2.4 keV band \\citep{article:BrielHenryBohringer:1992}. Observations with XMM-Newton \\citep{article:Briel_etal:2001} revealed substructure in the X-ray halo supported by substantial turbulent pressure of at least $\\sim 10 \\%$ of the total pressure \\citep{article:Schuecker_etal:2004}. The Coma cluster is a natural candidate for gamma-ray observations.  In this article, results from the VERITAS observations of the Coma cluster of galaxies are reported, with complementing analysis of available data from the Large Area Telescope (LAT) on board the \\Fermi Gamma-ray Space Telescope. The VERITAS and \\Fermi-LAT data have been used to place constraints on cosmic-ray particle populations, magnetic fields, and dark matter in the cluster. Throughout the analyses, a present day Hubble constant of $H_{0} = 100h$ km s$^{-1}$ Mpc$^{-1}$ with $h=0.7$ has been used. ",
        "conclusions": "We have reported on the observations of the Coma cluster of galaxies in VHE gamma rays with VERITAS and complementary observations with the \\emph{Fermi}-LAT. VERITAS observed the Coma cluster of galaxies for a total of 18.6 hours of high-quality live time between March and May in 2008. No significant excess of gamma rays was detected above an energy threshold of $\\sim 220$ GeV. The \\emph{Fermi}-LAT has observed the Coma cluster in all-sky survey mode since its launch in June 2008. We have used all data available from launch to April 2012 for an updated analysis compared to published results \\citep{article:Ackermann_etal:2010}. Again, no significant excess of gamma rays was detected. We have used the VERITAS and \\emph{Fermi}-LAT data to calculate flux upper limits at the 99\\% confidence level for the cluster core (considered as both a point-like source and a spatially-extended emission region) and for three member galaxies. The flux upper limits obtained were then used to constrain properties of the cluster.  We have employed various approaches to constrain the CR population and magnetic field distribution that are complementary in their assumptions and hence well suited to assessment of the underlying Bayesian priors in the models. (1) We used a simplified ``isobaric CR model'' that is characterized by a constant CR-to-thermal pressure fraction and has a power-law momentum spectrum. While this model is not physically justified {\\em a priori}, it is simple and widely used in the literature and captures some aspects of more elaborate models such as (2) the simulation-based analytical approach of \\citet{article:PinzkePfrommer:2010}. The latter is a ``first-principle approach'' that predicts the CR distribution spectrally and spatially for a given set of assumptions. It is powerful since it only requires the density profile as input due to the approximate universality of the CR distribution (when neglecting CR diffusion and streaming). Note, however, that inclusion of these CR transport processes may be necessary to explain the radio-halo bimodality. (3) Finally, we used a pragmatic approach which models the CR and magnetic distributions in order to reproduce the observed emission profile of the Coma radio halo. While this approach is also not physically justified, it is powerful because it shows what the ``correct'' model has to achieve and can point in the direction of the relevant physics.  Within this pragmatic approach, we employ two different methods. Firstly, adopting a high magnetic field everywhere in the cluster ($B\\gg B_\\rmn{CMB}$) yields the minimum gamma-ray flux in the hadronic model of radio halos which we find to be a factor of 20 (60) below the most constraining flux upper limits of \\Fermi-LAT (VERITAS). Secondly, by matching the radio-emission profile (i.e., fixing the radial CR profile for a given magnetic field model) and by requiring the pion-decay gamma-ray flux to match the flux upper limits (i.e., fixing the normalization of the CR distribution), we obtain lower limits on the magnetic field distribution under consideration. Our limits for the central magnetic field range from $B_{0} = 0.5$ to $1.4\\,\\mu$G (for VERITAS flux limits) and from $B_{0} = 1.4$ to $5.5\\,\\mu$G (for \\Fermi-LAT flux limits). Since all {\\em (but one) of} these lower limits on $B_0$ are below the values favored by Faraday RM, $B_{0} = 4.7^{+0.7}_{-0.8}\\,\\mu$G \\citep{article:Bonafede_etal:2010}, the hadronic model is a very attractive explanation of the Coma radio halo.  {The \\Fermi-LAT upper limits start to rule out the parameter combination of $\\alpha_{B}\\gtrsim 0.7$ and $\\alpha \\gtrsim 2.5$ for the hadronic model of the Coma radio halo.}  Applying our simplified ``isobaric CR model'' to the most constraining VERITAS limits, we can constrain the CR-to-thermal pressure ratio, $X_\\CR$, to be below 0.048--0.43 (for a CR or gamma-ray spectral index, $\\alpha$, varying between 2.1 and 2.5). We obtain a constraint of $X_\\CR<0.1$ for $\\alpha=2.3$, the spectral index predicted by simulations at energies around 220 GeV. This limit is more constraining by a factor of 1.6 than that of the simulation-based model which gives $X_\\CR<0.16$. This difference is due to the concave form of the simulated spectrum which provides more pressure at GeV energies in comparison to a pure power-law spectrum of $\\alpha=2.3$.  The \\Fermi-LAT flux limits constrain $X_\\CR$ to be below 0.012--0.017 (for $\\alpha$ varying between 2.3 and 2.1), only weakly depending on the assumed CR spectral index. Assuming $\\alpha=2.1$, which is very close to the simulated spectral index for the energy range of 1--3 GeV, we obtain a constraint which is identical to that from our simulation-based model within the virial radius of $X_\\CR<0.017$. That constraint improves to $X_{\\CR}<0.012$ for an aperture of 0.4$^\\circ$ corresponding to a physical scale of $R \\simeq R_{200}/3 \\simeq 660$~kpc.  {These upper limits are now starting to constrain the CR physics in self-consistent cosmological cluster simulations and cap the maximum CR acceleration efficiency at structure formation shocks to be $<50\\%$. Alternatively, this may argue for non-negligible CR transport processes such as CR streaming and diffusion into the outer cluster regions \\citep{article:Aleksic_etal:2012}.} These are encouraging results in that we constrain the CR pressure (of a phase that is fully mixed with the ICM) to be at most a small fraction ($<0.017$) of the overall pressure. As a result, hydrostatic cluster masses and the total Comptonization parameter due to the Sunyaev-Zel'dovich effect suffer at most a very small bias due to CRs.  We have also used the flux upper limits obtained with VERITAS to constrain the thermally-averaged product of the total self-annihilation cross section and the relative velocity of DM particles. Modeling the Coma cluster DM halo with a NFW profile we derived limits on $\\expval{\\sigma v}$ to be on the order of $10^{-20}$ to $10^{-21}$ cm$^{-3}$ s$^{-1}$ depending on the chosen aperture. These limits are based on conservative estimates of the astrophysical factor, where a possible boost to the annihilation rate due to DM substructures in the cluster halo has been neglected. Including such a boost could scale down the present limits by a factor $O(1000)$ in the most optimistic cases."
    },
    "1208/1208.1785_arXiv.txt": {
        "abstract": "We present results of observations in the UV to near-IR range for eight blazars, three of which have been recently discovered at Very High Energies (VHE) and five appearing as interesting candidates for VHE $\\gamma$-ray detection. We focus in this paper on the search for their redshifts, which are unknown or considered as uncertain. ",
        "introduction": "The redshift determination for blazars is often difficult as the bright non-thermal emission of the jet easily hides the host galaxy. For blazars detected in the VHE range, this measurement is very important as it allows to model the imprint of the Extragalactic Background Light (EBL) on their emission constraining models of both. ",
        "conclusions": "We demonstrate that the lower limit proposed for PKS~0447--439 at $z\\ge$1.246 by \\cite{landt_2012} is not correct, and find no evidence of the feature used by \\cite{perlman_1998} to propose $z=$0.205. We place a clear lower limit at z$\\geq$0.506 to the redshift of KUV~00311--1938, but do not see evidence for the feature used by \\cite{pirano_2007} to propose the tentative redshift $z=$0.61. We confirm the redshift proposed by \\cite{falomo_2000} for PKS~0301--243 and improve its precision at $z=$0.266. Finally, for a set of five blazars considered as good candidates for TeV emission, we give three firmly determined redshifts, a tentative redshift and a lower limit. We encourage their observation by current VHE $\\gamma$-ray observatories.  \\begin{theacknowledgments} S. Pita wishes to thank the ESO staff and in particular C. Martayan for their help in performing the observations. \\end{theacknowledgments}"
    },
    "1208/1208.3055_arXiv.txt": {
        "abstract": "An attractive mechanism for radiation transport by electron holes from the magnetospheric auroral cavity source region down to the ionosphere and possibly even down to the atmosphere is examined. Because of the excitation and propagation properties of the X mode, this mechanism turns out to be highly improbable for the usual resonant excitation of radiation of frequency just below the local electron cyclotron frequency, $\\omega\\lesssim\\omega_{ce}$. It could work only, if the auroral ionosphere would be locally perforated being of sufficiently low density for allowing electron holes  riding down to low altitudes on the auroral electron beam. If resonant excitation of higher electron cyclotron harmonics, $\\omega\\sim \\ell\\omega_{ce}, ~\\ell= 2,3,\\dots$, becomes possible, a still unexplored mechanism, then radiation excited inside the hole could be transported to lower altitudes than generated. If this happens, radiation would provide another mechanism of coupling between the magnetospheric plasma and the atmosphere. ",
        "introduction": "This Letter shows that the otherwise very attractive idea of transport of narrow band auroral kilometric radiation generated inside electron holes \\citep[the generation mechanism has been elucidated in][]{treumann2011a,treumann2011b} in the auroral cavity  down to the ionosphere and possibly even further down to the neutral atmosphere is (probably) quite unrealistic.  Such a radiation transport might occur just under extremely favorable circumstances prevalent in the ionosphere when the ionosphere becomes perforated. Unfortunately, under any other or normal conditions, however, electromagnetic radiation transport by holes, as can naively be guessed, will probably not take place.  We do not explore the transport mechanism in detail, for a detailed calculation requires extended access to the space-time distribution of the electron densities and temperatures at altitudes reaching from the atmosphere up to the magnetosphere, all longitudes and latitudes in the interval from subauroral to polar latitudes. This effort is outside the scope of the present communication. However, we feel that the mere idea is interesting and tempting enough for excluding it as a possible radiative coupling mechanism between the atmosphere and magnetosphere from upside down, at least at frequencies below $\\omega_{ce}$.  Coupling between the atmospheric and plasma environments of Earth has recently become of increasing interest. There is no doubt that such coupling really exists. It is, however, believed that it is mainly upward and of basically mechanical and thermal nature. The various known coupling processes have recently been reviewed \\citep{bosinger2012} to some detail. Some surprising effects of Earth's geomagnetic field on the atmospheric density distribution have been discovered by the CHAMP spacecraft \\citep{luhr2012}.  Otherwise the effect of the magnetospheric plasma on the atmosphere is probably restricted mainly to auroral and polar latitudes and depends heavily on the magnetospheric particle distribution and its space-time and energy-time variability.  On the other hand, from a magnetospheric point of view,  some violent processes seem not to have any known or remarkable effect in the ionosphere and atmosphere. Such processes relate to the generation of auroral kilometric radiation which is known to contain several per cent of the total energy content of a magnetospheric substorm, which in radiation is a huge percentage. Though this number is still energetically negligible in any large-scale atmospheric or ionospheric process, the question arises whether it could be transported at all down to the ionosphere and atmosphere where it would be detected as banded free-space radiation. Interestingly, \\citet{labelle2011b} recently reported the probable observation of auroral kilometric radiation even at the ground from Antarctica. In the upper ionosphere banded (so-called MF burst) radiation has also been observed though it is not yet finally decided whether it is electromagnetic or electrostatic \\citep{labelle2011a}.\\footnote{For plasma wave observations see \\citet{labelle1988,labelle2002}.} Its generation mechanism remains still unclear as also its relation to substorms. Here, we take these just as examples while not claiming that they could be explained by the mechanism we are going to check below. ",
        "conclusions": ""
    },
    "1208/1208.4443_arXiv.txt": {
        "abstract": "New spectroscopic observations show that the double degenerate system NLTT~16249 is in a close orbit ($a=5.6\\pm0.3\\,R_\\odot$) with a period of 1.17~d. The total mass of the system is estimated between 1.47 and 2.04\\,$M_\\odot$ but it is not expected to merge within a Hubble time-scale ($t_{\\rm merge}\\approx 10^{11}$ yr). Vennes \\& Kawka (2012, ApJ, 745, L12) originally identified the system because of the peculiar composite hydrogen (DA class) and molecular (C$_2$--DQ class--and CN) spectra and the new observations establish this system as the first DA plus DQ close double degenerate. Also, the DQ component was the first of its class to show nitrogen dredged-up from the core in its atmosphere. The star may be viewed as the first known DQ descendant of the born-again PG1159 stars. Alternatively, the presence of nitrogen may be the result of past interactions and truncated evolution in a close binary system. ",
        "introduction": "The high-proper motion star NLTT~16249 is a double degenerate system showing hydrogen lines and molecular carbon and cyanogen bands in its optical spectrum \\citep{ven2012}. Although large radial velocity variations were noted by \\citet{ven2012}, the orbital parameters, period and separation, are yet to be determined. The detection of photospheric nitrogen in the carbon-rich (DQ) component of this system was a first occurrence for this class of objects and, according to the dredge-up scenario commonly applied to DQ white dwarfs \\citep{pel1986,mac1998}, it implied the presence of nitrogen in the white dwarf core. The other component is a hydrogen-rich (DA) white dwarf with a mass above average but with a luminosity similar to that of the DQ white dwarf.  The two characteristics of the system, a nitrogen-enriched DQ component and a likely close orbit, may or may not be related. \\citet{alt2005} established a clear evolutionary link between born-again stars and a nitrogen/oxygen enrichment in PG1159 stars and their DQ descendents. Members of the PG1159 class have a helium-rich surface with notable enrichment in carbon, nitrogen and oxygen \\citep{wer1991,dre1998,wer2006}. \\citet{alt2005} argued that not all DQ stars followed the born-again path because the carbon abundance is lower in some of these objects than would be expected following this evolutionary path. However, the $\\log{\\rm C/He}$ versus \\teff\\ trend established by \\citet{duf2005} and \\citet{koe2006} does appear to follow the model predictions of \\citet{alt2005}. Moreover, \\citet{duf2005} noted two discernible tracks with one at a higher carbon abundance that was attributed to a thinning of the helium layer. Therefore, data and born-again models may be compatible for the bulk of DQ white dwarfs. However, nitrogen is not detected in any DQ stars other than in the DQ in NLTT~16249.  In this Letter, we present convincing evidence that the components of the binary NLTT~16249 are in a close orbit. Our new radial velocity measurements (Section 2) and our original model atmosphere analysis \\citep{ven2012} help constrain the component properties and offer clues to the origin and evolutionary prospect of the system (Section 3). We summarize and discuss these new results in Section 4. ",
        "conclusions": "We have shown that the peculiar DQ white dwarf in NLTT~16249 is in a close 1.17 d orbit with a DA white dwarf companion. Close double-degenerate stars are common and \\citet{max1999} reported a fraction of 5 to 19\\% of close pairs in their radial velocity survey. White dwarfs with composite spectra indicative of DA plus DB (He{\\sc i}) or DA plus DC (cool He-rich) pairs are also relatively common in large surveys \\citep[see, e.g.,][]{lim2010,tre2011}. However, only two DA+DQ pairs are known: NLTT~16249 and SDSS~J153210.04+135615.0 (NLTT~40489) which was recently identified by \\citet{gia2012}.  The velocity amplitudes in NLTT~16249 imply that the DQ is slightly more massive than its companion in which case it is more likely that the DQ formed first, both from $\\sim3\\,M_\\odot$ progenitors, unless stable mass transfer reversed the initial mass ratio \\citep[for detailed scenarios see][]{nel2001}. The relatively long orbital period reported here rules out the prospect of a merger within a Hubble time.  The DQ white dwarf is peculiar and the nitrogen concentration in its atmosphere is, so far, unique. \\citet{alt2005} followed the evolution of a 2.7$M_\\odot$ main-sequence star past the PG1159 spectroscopic stage and down the cooling track up to the DQ stage (0.6$M_\\odot$). The star follows a ``born-again'' He-flash loop that resets its evolution onto the post-asymptotic giant-branch (post-AGB), considerably modifying the chemical structure of the outer layers that an otherwise normal post-AGB would possess. In particular, \\citet{alt2005} predicted measurable concentrations of nitrogen and oxygen in the atmosphere of DQ descendents of born-again stars.  The nitrogen abundance measured in the DQ component of NLTT~16249 is qualitatively similar to the abundance predicted for the 10,500~K model, the lowest temperature considered by \\citet{alt2005}. However, the ${\\rm C/N}$ ratio for that particular model is close to 300, i.e, much larger than measured in NLTT~16249 (${\\rm C/N}\\approx 50$). In the born-again context, the difference could be attributed to a different progenitor mass which could affect the outcome of the chemical evolution, although such models are not available to us. \\citet{alt2005} also predicted a concentration of oxygen between that of carbon and nitrogen (${\\rm C/O}\\approx 100$). It is possible  to constrain the ${\\rm C/O}$ abundance ratio in DQ white dwarfs using the near and far ultraviolet CO ``Fourth-Positive\" (4P) bands along with CN violet and C$_2$ Swan bands. Finally, the $^{13}{\\rm C}/^{12}{\\rm C}$ isotopic ratio expected for the same evolutionary track is well below detection limits reported by \\citet{ven2012} in the case of the DQ in NLTT~16249.  What would be the role, if any, played by binary interactions in shaping the chemical structure of the DQ progenitor in NLTT~16249 ? The DQ white dwarfs are not known to show nitrogen in their spectra \\citep[see][]{duf2005}. Therefore, it may not be a coincidence that the only known detection of nitrogen in a DQ white dwarf is also that of a white dwarf in a close double-degenerate system. In that context as well, it remains to be shown that detectable traces of nitrogen would be the expected outcome of mass-transfer or early envelope ejection."
    },
    "1208/1208.6446_arXiv.txt": {
        "abstract": "{The study of Type Ia supernovae (SNIa) has lead to greatly improved insights into many fields in astrophysics, e.g. cosmology, and also into the metal enrichment of the universe. Although a theoretical explanation of the origin of these events is still lacking, there is a general consensus that SNIa are caused by the thermonuclear explosions of carbon/oxygen white dwarfs with masses near the Chandrasekhar mass. } {We investigate the potential contribution to the supernova Type Ia rate from the population of merging double carbon-oxygen white dwarfs. We aim to develope a model that fits the observed SNIa progenitors as well as the observed close double white dwarf population. We differentiate between two scenarios for the common envelope (CE) evolution; the $\\alpha$-formalism based on the energy equation and the $\\gamma$-formalism that is based on the angular momentum equation. In one model we apply the $\\alpha$-formalism always. In the second model the $\\gamma$-formalism is applied, unless the binary contains a compact object or the CE is triggered by a tidal instability for which the $\\alpha$-formalism is used. } {The binary population synthesis code SeBa was used to evolve binary systems from the zero-age main sequence to the formation of double white dwarfs and subsequent mergers. SeBa has been thoroughly updated since the last publication of the content of the code. } { The limited sample of observed double white dwarfs is better represented by the simulated population using the $\\gamma$-formalism for the first CE phase than the $\\alpha$-formalism. For both CE formalisms, we find that although the morphology of the simulated delay time distribution matches that of the observations within the errors, the normalisation and time-integrated rate per stellar mass are a factor $\\sim 7-12$ lower than observed. Furthermore, the characteristics of the simulated populations of merging double carbon-oxygen white dwarfs are discussed and put in the context of alternative SNIa models for merging double white dwarfs. } {} ",
        "introduction": "Type Ia supernovae (SNIa) are one of the most energetic explosive events known. They have been of great importance in many fields, most notably as a tool in observational cosmology. They have been used very successfully as standard candles on cosmological distance scales \\citep[e.g.][]{Rie98, Per99}, owing to the special property of great uniformity in the light curves \\citep[e.g.][]{Phi93}. The SNIa also strongly affect the Galactic chemical evolution through the expulsion of iron \\cite[e.g.][]{Ber91}.  Despite their significance Type Ia supernovae are still poorly understood theoretically.  Supernovae Type Ia are generally thought to be caused by thermonuclear explosions of carbon/oxygen (CO) white dwarfs (WDs) with masses near the Chandrasekhar mass $M_{\\rm ch}\\approx 1.4\\Mo$ \\citep[e.g.][]{Nom82}. Various progenitor scenarios have been proposed. The standard scenarios can be divided into two schools of thoughts: the single-degenerate (SD) \\citep{Whe73} and double-degenerate (DD) scenario \\citep{Web84, Ibe84}. In the SD scenario, a CO WD explodes as an SNIa if its mass approaches $M_{\\rm ch}$ through accretion from a non-degenerate companion. In the DD scenario, two CO WDs can produce an SN Ia while merging if their combined mass is larger than $M_{\\rm ch}$.  However, observationally as well as theoretically, the exact nature of the SNIa progenitors remains unclear. The explosion mechanism is complex due to the interaction of hydrodynamics and nuclear reactions. Several models exist that vary for example between a detonation or deflagration disruption or vary between explosions at the Chandrasekhar mass or sub-Chandrasekhar masses \\citep[see e.g.][for a review]{Hil00}. It also remains unclear whether the DD and SD scenario both contribute to the SNIa rate or if one of the scenarios dominates. Both scenarios have problems in matching theories with observations. A serious concern about the DD scenario is whether the collapse of the remnant would lead to a supernova or to a neutron star through accretion-induced collapse \\citep[see][]{Nom85, Sai85, Pie03, Yoo07, Pak10, Pak12, She12}. Although in the SD channel the models for the explosion process need to be fine-tuned to reproduce the observed spectra and light curves, an SNIa like event is more easily reproduced in the simulations of the explosion process. One problem with the SD scenario is that the white dwarfs should go through a long phase of supersoft X-ray emission, although it is unclear if there are enough of these sources to account for the SNIa rate \\citep[see][]{DiS10, Gil10, Hac10}. Moreover archival data of known SNIa have not shown this emission unambiguously, but there is may be one case \\citep[see][]{Vos08, Roe08, Nie10}. Furthermore, SNIa that take place more than a few 10$^9$ years after the starburst \\citep[see e.g.][]{Mao10} are hard to create in this channel \\citep[e.g.][]{Yun00, Han04}.  To use SNIa as proper standard candles, we need to know what SNIa are, when they happen and what their progenitors are. Therefore, we study the binary evolution of low- and intermediate mass stars. In a forthcoming paper (Bours, Toonen \\& Nelemans, in preparation) we study the SD-scenario by looking into the poorly understood physics of accretion onto white dwarfs. In this paper we focus on the DD scenario and the effect of the as yet very uncertain phases of common envelope (CE) evolution on the double white dwarf (DWD) population. These DWD systems are interesting sources for studying various phases of stellar evolution, in our case the CE evolution. Gravitational wave emission is also important as this affects the binary system by decreasing the orbital period and eventually leading to a merger \\citep{Kra62, Pet64}, or possible a SNIa. The DWDs are expected to be the dominant source \\citep{Eva87, Nel01c} of gravitational waves for the future space-born gravitational wave observatories such as eLISA \\citep{Ama12, Ama12b}.   We study the population of merging DWDs that might lead to a SNIa from a theoretical point of view. We incorporated results from observations where possible. We use the population synthesis code \\SeBa\\ for simulating the stellar and binary evolution of stellar systems that leads to close DWDs. In Sect. \\ref{sec:seba} we describe the code and the updates since the last publication of \\SeBa. A major influence on the merging double-degenerate population is the poorly understood CE phase \\citep{Pac76, Web84, Nel00}. We adopt two different models for the CE. In Sect. \\ref{sec:ce} we describe these models and their implications for the observations of close DWDs. In Sect. \\ref{sec:path} we discuss the binary paths leading to SNIa for each model. The SNIa rates and time-integrated numbers are derived in Sect. \\ref{sec:dtd}. The properties of the population of merging DWDs are discussed in the context of the classical and  alternative sub- and super-Chandrasekhar SNIa explosion models in Sect. \\ref{sec:pop}. A discussion and conclusion follows in Sect. \\ref{sec:con}.   \\begin{figure*} \\centering \\begin{tabular}{c c} \\includegraphics[angle=270, scale = 0.4]{EFT_ag_M1.ps} & \\includegraphics[angle=270, scale = 0.4]{SeBa_r80_ag_M1.ps} \\\\ (a) & (b) \\\\ \\end{tabular} \\caption{Simulated population of visible double white dwarfs as a function of orbital period and mass of the brighter white dwarf. Left: the stellar evolution tracks according to EFT are used; right: HPT (using model $\\mga$, see Sect. \\ref{sec:ce}). The intensity of the grey scale corresponds to the density of objects on a linear scale. The same grey scale is used for both plots. Observed binary white dwarfs are overplotted with filled circles. Thick points taken are from \\citet{Mar11}, thinner points from \\citet{Tov04, Nap05, Kul10, Bro10, Bro11, Mar11, Kil11, Kil11b, Kil11c}, see Sect. \\ref{sec:seba_obs} for a discussion.} \\label{fig:pop_M1} \\end{figure*} ",
        "conclusions": "\\label{sec:con} We studied the population of SNIa progenitors from merging double CO WDs with a combined mass exceeding the Chandrasekhar mass, the so-called DD progenitors. We considered two prescriptions of the CE phase. The CE evolution is a crucial ingredient in the formation of close double degenerate compact objects, but the process itself is still poorly understood. The first model assumes the $\\alpha$-formalism for all CE. The second model is a combination of the $\\alpha$-formalism and the $\\gamma$-formalism (see Sect. \\ref{sec:ce}). Typically, the first CE is described by the $\\gamma$-scenario and the second by the $\\alpha$-formalism, if mass transfer is unstable.  We applied the updated version of the population synthesis code SeBa to simulate the population of DWDs and SNIa progenitors. At present, close DWDs (of all WD types) are the closest related systems to the DD SNIa progenitors that are visible in bulk. The mass ratio distribution of the DWDs in model $\\maa$ is inconsistent with the observations. Using model $\\mga$ the simulated population of DWDs compares well with observations, nevertheless, this is what the $\\gamma$-formalism was designed to do.  Recently, \\citet{Web08} and \\citet{Zor10} claimed that the predictive power of the $\\gamma$-scenario is more restricted. They suggested that the $\\alpha$-scenario is valid when sources of an energy other than the binding energy of the envelope is available, such as, the energy released by recombination in the common envelope. This could explain the high value of $\\alpha$ found by \\citet{Nel00} for the second CE, but certainly does not solve the problem for the first CE for which \\citet{Nel00} found a value of $\\alpha < 0$.  The delay time distributions from our two models show the characteristic shape of a strong decay with time. This strong decay is expected when the delay time is dominated by the gravitational wave timescale ($t_{gr} \\propto a^4$) and the distribution of orbital separations at DWD formation is similar to the initial (ZAMS) distribution of $N(a) da \\propto a^{-1} da$ \\citep{Abt83}. The DTD from model $\\mga$ fits the shape of the observed DTD best. \\citet{Men10} also showed a DTD using the $\\gamma$-scenario for the CE phase. They found that the DD DTD lies almost an order of magnitude lower in absolute rate than when using the $\\alpha$-scenario. However, they used the $\\gamma$-formalism for all CE phases. In our prescription (see Sect. \\ref{sec:ce}) the $\\gamma$-formalism is typically used in the first CE phase only. The reason for this is that in equal mass systems there is more angular momentum compared to unequal mass systems with similar orbits. \\citet{Men10} and also \\citet{Yun00, Rui09} and \\citet{Cla11} showed DD DTDs using the $\\alpha$-scenario (however their CE-prescriptions may differ slightly from Eq.\\ref{eq:Egr}). Surprisingly, but as realised before, even though different groups used different binary evolution codes with different versions of the $\\alpha$-CE and CE efficiencies, the DTDs of the DD channel are very uniform in that they show a strong decline with time \\citep[see for example][for an overview]{Nel12}.  Usually in synthetic DTD studies, the shape and normalisation of the DTD are discussed separately. This might not be valid any more, as more and more observed rates are available and the conversion from observational units to synthetic units (e.g. the star formation history (SFH) and rate in per K-band luminosity instead of per \\Msolar\\ of created stars) is better understood. For example, the SFH is often convolved with the DTD to estimate the SNIa rate in spiral galaxies like our Milky Way. The problem with this is that different assumptions for the Galactic SFH can significantly alter the theoretical Galactic SNIa rate. Since the SNIa rate follows the SFH with typical delay times of a few Gyr, the synthetic Galactic SNIa rate is very sensitive to the assumed SFH at recent times. When a constant SFH (of $\\gtrsim 3\\Mo\\peryr$) is assumed, the SNIa rate is artificially enhanced compared with detailed SFHs that show a peak in the star formation at a few Gyr and a decline to $1\\Mo\\peryr$ at recent times, see e.g. \\citet{Nel04}. In the observed SNIa rates of \\citet{Mao10b} and \\citet{Mao11} the detailed SFH of every individual galaxy or galaxy subunit was taken into account to reconstruct the DTD. Therefore it is no longer necessary to convolve the theoretical SNIa rate from a burst of star formation with an approximate SFH. The theoretical calculations of the SNIa rate from a single starburst can directly be compared with observations.  We found that the normalisation of the DTD of model $\\maa$ and $\\mga$ do not differ much, even though the CE evolution is very different. The time-integrated number of SNIa in model $\\maa$ ($3.3 \\cdot 10^{-4} \\Mo^{-1}$) is 70\\% larger as in model $\\mga$ ($2.0\\cdot 10^{-4} \\Mo^{-1}$). But most importantly, the simulated time-integrated numbers do not match the observed number of $2.3\\pm 0.6 \\cdot 10^{-3} \\Mo^{-1}$ by \\citep{Mao11} by a factor of $\\sim 7-12$. If our understanding of binary evolution and initial parameter distributions is correct, the standard DD channel is not a major contributor to the SNIa rate.  For the SNIa model proposed by \\citet{Pak10}, in which carbon burning is ignited in the merger process of two massive white dwarfs of nearly equal mass, we found an SNIa rate of $2.3 \\cdot 10^{-5} \\Mo^{-1}$ for model $\\mga$ and $1.8 \\cdot 10^{-5} \\Mo^{-1}$ for model $\\maa$. \\citet{Pak10} founds that these systems resemble subluminous SNIa such as SN 1991bg. Assuming the fraction of 1991bg-like events to all SNIa events is 15$\\pm$6\\% \\citep{Li01, Li11}, the observed event rate is $3.5 \\pm 1.6 \\cdot 10^{-4} \\Mo^{-1}$. \\citet{Ker10} proposed a model in which sub-Chandrasekhar WDs can explode as an SNIa. In this scenario two white dwarfs of nearly equal mass merge, though carbon ignition occurs only after the merger when the thick disk surrounding the remnant is accreted onto it. The event rate is $2.3 \\cdot 10^{-4} \\Mo^{-1}$ for model $\\mga$ and $1.9 \\cdot 10^{-4} \\Mo^{-1}$ for model $\\maa$. When only taking into account systems with a combined mass below 1.4$\\Mo$, the rates are $1.3 \\cdot 10^{-4} \\Mo^{-1}$ and $8.8 \\cdot 10^{-5} \\Mo^{-1}$, respectively. In the scenario proposed by \\citet{Pak10} and in the scenario by \\citet{Ker10}, systems are required to have high mass ratios. We showed that the mass ratio distribution of DWDs depends on the prescription for the CE. When the $\\gamma$-scenario is used, the average mass ratio of DWDs lies closer to one, which increases the SNIa rate in the above described scenario with respect to the $\\alpha$-scenario. The rates of the channel proposed by \\citet{Ker10} for systems with sub-Chandrasekhar masses are on the same order of magnitude as the rates of the standard DWD channel. Therefore the combination of the two models is not sufficient to explain the observed rates. For our synthetic rates of the DD scenario to match the observed SNIa rates, within our current model of binary evolution, the parameter space of the DD progenitors has to be increased severely, e.g. to include all CO-CO and CO-He mergers, which seems unlikely.  Alternatively (and if contributions from channels other than the DD are minor), our model underpredicts the fraction of standard DD SNIa progenitors in the entire DWD population. Our model of the visible population of DWDs predicts 0.9-2.9\\% of the visible DWDs (depending on the model) to be SNIa progenitors. To match the observed rate of \\citet{Mao11}, 10-30\\% (excluding any errors on the observed and synthetic values) of the observed DWDs should lie in the SNIa progenitor region (upper left corner of Fig. \\ref{fig:pop_Mt}). With 46 observed DWDs so far, 4-15 SNIa progenitors are expected without taking non-uniform selection effects into account. So far, only two systems have been found that possibly are SNIa progenitors, which makes it improbable, but not impossible, that our model underpredicts the number of DD SNIa progenitors. When the population of observed DWDs is increased, the fraction of SNIa progenitors amongst DWDs will give more insight into the validity of our knowledge of binary evolution of massive DWDs.  Concluding, although the shape of the DD DTD fits the observed DTD beautifully, the normalisation does not. An important point is that we did not optimise our model to fit the observed DTD in shape or number. We showed that the normalisation can be influenced by the metallicity; $\\sim 30-60\\%$ depending on the model for $Z=0.001$ with respect to $Z=0.02$. Furthermore, the normalisation depends on the initial mass function, the percentage of single stars, and the initial distribution of mass ratios and orbital periods. In this paper and in \\citet{Nel12} we assumed the percentage of single stars to be 50\\%. Results from e.g. \\citet{Kou07} and \\citet{Rag10} showed that the binary fraction might be as high as 70\\% or more for A- and B-type stars, potentially raising the synthetic SNIa rate by a factor $<2$. Preliminary results show that the initial distribution of mass ratio and orbital separation affects the slope of the DTD, still the strong decline with time remains. Moreover, the integrated rates are not affected by factors sufficient to match the observed rate. Additional research is needed to study if the normalisation can be raised sufficiently to match the observed rate. If not, the main contribution to the SNIa rate comes from other channels, such as the SD scenario (e.g. supersoft sources), double detonating sub-Chandrasekhar accretors \\citep[see e.g.][]{Kro10}, or Kozai oscillations in triple systems (\\citealt{Sha12}; Hamers et al. in prep.)."
    },
    "1208/1208.1320_arXiv.txt": {
        "abstract": "We use an extended version of the hadronic SU(3) non-linear realization of the sigma model that also includes quarks to study hybrid stars. Within this approach, the degrees of freedom change naturally as the temperature/density increases. Different prescriptions of charge neutrality, local and global, are tested and the influence of strong magnetic fields and the anomalous magnetic moment on the particle population is discussed. ",
        "introduction": "In order to describe hybrid stars we make use of a model that contains both hadronic and quark degrees of freedom. At high densities and/or temperatures the hadrons are replaced by quarks, as their effective masses increase (hadrons) and decrease (quarks). The aforementioned model (see \\cite{Dexheimer:2009hi} for details) is an extended version of the SU(3) non-linear realization of the sigma model \\cite{Papazoglou:1998vr,Dexheimer:2008ax,Negreiros:2010hk} within the mean-field approximation. Changes in the order parameters $\\sigma$ and $\\Phi$ signal chiral symmetry restoration and quark deconfinement, respectively. The potential for $\\Phi$ is an extension of the Polyakov loop potential \\cite{Roessner:2006xn} modified to also depend on baryon chemical potential. In this way the model is able to describe the entire QCD phase diagram and agrees with lattice QCD constraints \\cite{Fodor:2004nz}. The phase transitions at low temperatures and high densities are of first order, while at high temperatures and low densities the model exhibits a smooth crossover, in accordance with lattice QCD results..  For neutron star calculations, matter is further required to be charge neutral and in chemical equilibrium. In the case of a first-order phase transition, we can either require each phase to be individually charge neutral (local charge neutrality) or allow both phases to be charge neutral only when combined (global charge neutrality). Fig.~\\ref{plot1} shows how the resulting QCD phase diagram looks like following both of these prescriptions. When global charge neutrality is assumed, a mixed phase appears. It extends over a range of baryon chemical potentials, that corresponds to a range of baryon densities (Fig.~\\ref{plot2}). If local charge neutrality is assumed instead, the phase transition occurs for a given temperature at a single value of the chemical potential. Nevertheless, a small step in baryon chemical potential at the phase transition corresponds to a large jump in baryon density. Note that for local charge neutrality all densities within these limits correspond to the same pressure and, therefore, cannot occupy a physical space in a star.  \\begin{figure} \\includegraphics[height=.3\\textheight]{plot1} \\caption{Temperature versus baryon chemical potential for neutron star matter assuming local and global charge neutrality.} \\label{plot1} \\end{figure}  \\begin{figure} \\includegraphics[height=.3\\textheight,clip,trim=0 0 0 1]{plot2} \\caption{Temperature versus baryon density for neutron star matter assuming local and global charge neutrality.} \\label{plot2} \\end{figure}  The next step towards a complete description of neutron stars is the inclusion of the effect of magnetic fields in the model. The highest magnetic field observed on the surface of magnetars is on the order of $10^{15}$ G. The highest possible magnetic field in the center of stars, on the other hand, can only be estimated through models, even when applying the Virial theorem. Some results indicate limiting magnetic fields ranging between $B=10^{18}-10^{20}$ G ~\\cite{Bocquet:1995je,Cardall:2000bs,1991ApJ...383..745L,Chakrabarty:1997ef,Bandyopadhyay:1997kh,Broderick:2001qw,Ferrer:2010wz, Malheiro:2004sb}. In order to avoid hydrodynamical instabilities due to pressure anysotropy \\cite{Chaichian:1999gd,PerezMartinez:2005av,PerezMartinez:2007kw,Huang:2009ue,Paulucci:2010uj} and still be able to include strong magnetic fields, we consider in our calculation a magnetic field in the z-direction that varies with baryon chemical potential. More precisely the assumed magnetic field $B^*$ runs from a surface value $B_{surf}=69.25$ MeV$^2=10^{15}$ G (when $\\mu_B=938$ MeV) to different central values $B_c$ at extremely high baryonic chemical potential following \\cite{Dexheimer:2011pz} \\begin{equation} B^*(\\mu_B)=B_{surf}+B_c[1-e^{b\\frac{(\\mu_B-938)^a}{938}}], \\end{equation} with $a=2.5$, $b=-4.08\\times10^{-4}$ and $\\mu_B$ given in MeV. As can be seen in Fig.~\\ref{Beff}, the values of effective magnetic field only approach $B_c$ at very high baryonic chemical potentials and, in practice, only about $70\\%$ of $B_c$ can be reached inside neutron stars. The use of baryon chemical potential instead of density was chosen to prevent discontinuities in the magnetic field at the phase transition where, as we have already mentioned, there is a jump in baryon density. Nevertheless, the constants $a$ and $b$ and the form of the $B^*$ expression were chosen to reproduce (in the absence of quarks) the effective magnetic field curve as a function of density from Refs.~\\cite{Bandyopadhyay:1997kh,Mao:2001fq,Rabhi:2009ih}.  \\begin{figure} \\includegraphics[height=.3\\textheight]{Beff} \\caption{Effective magnetic field as a function of baryonic chemical potential shown for different central magnetic fields. Recall that using Gaussian natural units $1$ MeV$^2=1.44\\times10^{13}$ G. \\label{Beff}} \\end{figure} ",
        "conclusions": "The magnetic field in the z-direction forces the eigenstates in the x and y directions of the charged particles to be quantized into Landau levels. The energy levels of all baryons are further split due to the alignment/anti-alignment of their spin with the magnetic field (anomalous magnetic moment effect AMM). A complete analysis of these effects together with chiral symmetry restoration and quark deconfinement within the extended version of the SU(3) non-linear realization of the sigma model on neutron star observables can be found in \\cite{Dexheimer:2011pz}. Similar studies using different models can be found in Refs.~\\cite{Menezes:2009uc,Menezes:2008qt,2011PhRvC..83f5805A}.  In this work, we chose to focus on the analysis of magnetic field effects on the chemical composition of the neutron star, using global charge neutrality. More specifically, Fig.~\\ref{popB} shows the density of fermions (with quark number densities divided by 3) when a central magnetic field $B_c=5\\times10^5$ MeV$^2=7.22\\times10^{18}$ G with AMM is considered. In the mixed phase the hadrons disappear smoothly as the quarks appear. The hyperons, despite being included in the calculation, are suppressed by the appearance of quarks and only a very small amount of $\\Lambda$'s appear. The strange quarks appear after the other quarks and do not make substantial changes in the system. The \"wiggles\" in the charged particle densities mark when their Fermi energies cross the discrete threshold of a Landau level. The charged particle population is enhanced due to $B$, as their chemical potentials increase.  Although the AMM is known to make the EOS stiffer, it does not have a very significant effect in the particle population \\cite{Broderick:2001qw}. This fact can be easily understood in terms of polarization, when, instead of looking at the total particle density (sum of spin up and down particle densities) for each species, we look at the spin up/spin down particle densities separately. Fig.~\\ref{popspin} shows that some of these particles are enhanced while others are suppressed. This comes from the definition of Landau level $\\nu=l+\\frac{1}{2}-\\frac{s}{2} \\frac{{Q_e}_i}{|{Q_e}_i|}$. Particles with spin projection $s$ opposite to their charge sign ${Q_e}_i/{|{Q_e}_i|}$ cannot have $\\nu=0$, even when their orbital angular momentum $l$ is equal to zero, causing them to have a smaller density when summing over all levels. Particles with spin projection equal to their charge sign are enhanced. Note that the electrons in the quark phase are fully polarized.  For non-charged baryons, the ones with negative spin projection are suppressed due to the increase in their effective masses caused by the anomalous magnetic moment ${m_i}^* \\to {m_i}^* - s \\kappa_i B$ combined with negative $\\kappa_i$ (like neutrons and $\\Lambda$'s). The baryons with negative spin projection and positive $\\kappa_i$ (like the protons) are enhanced by the AMM. Particles with spin up have the exactly opposite behavior.  \\begin{figure} \\includegraphics[height=.3\\textheight]{popB} \\caption{Particle densities as a function of baryonic chemical potential assuming global charge neutrality for $B_c=5\\times10^5$ MeV$^2=7.22\\times10^{18}$ G including AMM. \\label{popB}} \\end{figure}  \\begin{figure} \\includegraphics[height=.3\\textheight]{popspin} \\caption{Same as Fig.~\\ref{popB} but differentiating particles by spins. Black, blue and red stand for negative spin projections, while green, purple and orange stand for positive spin projections, respectively.\\label{popspin}} \\end{figure}  We have shown in this work some possible effects of strong magnetic fields in hybrid stars. The presence of different hadronic and quark degrees of freedom makes the extended version of the SU(3) non-linear realization of the sigma model an ideal tool for such an analysis, in all of the different possible phases. In the future, we intend to calculate the influence of strong magnetic fields in chiral symmetry restoration and quark deconfinement for neutron star matter at finite temperatures. Such kind of analysis can help us have a better understanding of the QCD phase diagram.   \\begin{theacknowledgments}  V. D. acknowledges support from CNPq (National Counsel of Technological and Scientific Development - Brazil). M. H. acknowledges support from the High Performance and High Productivity Computing (HP2C) project, the Swiss National Science Foundation (SNF) and ENSAR/THEXO. M. H. is also grateful for support from CompStar, a research networking program of the ESF.  \\end{theacknowledgments}"
    },
    "1208/1208.0689.txt": {
        "abstract": "We present new splitting methods  designed for the numerical integration of near-integrable Hamiltonian systems, and in particular for planetary N-body problems, when one is interested in very accurate results over a large time span. We derive in a systematic way an independent set of necessary and sufficient conditions to be satisfied by the coefficients of splitting methods to achieve a prescribed order of accuracy. Splitting methods satisfying such (generalized) order conditions are appropriate in particular for the numerical simulation of the Solar System described in Jacobi coordinates. We show that, when using Poincar\\'e Heliocentric coordinates,  the same order of accuracy may be obtained by imposing an additional polynomial equation on the coefficients of the splitting method. We construct several splitting methods appropriate for each of the two sets of coordinates by solving the corresponding systems of polynomial equations and finding the optimal solutions. The experiments reported here indicate that the efficiency of our new schemes is clearly superior to previous integrators when high accuracy is required.  \\vspace*{0.6cm}  \\begin{description} \\item $^1$Instituto de Matem\\'atica Multidisciplinar, Universitat Polit\\`ecnica de Val\\`encia, E-46022  Valencia, Spain. \\item $^2$Institut de Matem\\`atiques i Aplicacions de Castell\\'o and Departament de Matem\\`atiques, Universitat Jaume I, E-12071 Castell\\'on, Spain. \\item $^3$Astronomie et Syst\\`emes Dynamiques, IMCCE-CNRS UMR8028, Observatoire de Paris, UPMC, 77 Av. Denfert-Rochereau, 75014 Paris, France. \\item $^4$Konputazio Zientziak eta A.A. saila, Informatika Fakultatea, EHU/UPV, E-20018, Donostia/San Sebasti\\'an, Spain. \\end{description} ",
        "introduction": "\\label{Introduction}    Symplectic integrators have several features that turn out to be particularly appropriate when integrating numerically for long times evolution problems in dynamical astronomy. They preserve by construction the symplectic structure of the original Hamiltonian problem, so that the numerical solution inherits the qualitative properties of the exact one \\cite{sanz-serna94nhp}.  In particular, by using backward error analysis, it is possible to prove that this numerical solution is in fact exponentially close to the exact solution of a modified Hamiltonian.  Moreover, although the energy is not conserved along the trajectory, the error introduced by a symplectic method of order $r$ used with constant step size $\\tau$ is of order $\\mathcal{O}(\\tau^r)$ for exponentially long time intervals under rather general assumptions, whereas the error in position typically grows linearly with time \\cite{hairer06gni}.  Assume that, as is often the case, the Hamiltonian function is of the form $H(q,p) = T(p) + U(q)$, where the potential energy $U(q)$ depends on positions and the kinetic energy $T(p)$ is a function of the conjugate momenta. Then the equations of motion corresponding to $T(p)$ are trivially solvable, and the same happens with $U(q)$. By composing the flows of these two special Hamiltonian systems one gets a symplectic first order approximation to the exact flow. This simple composition constitutes an example of a symplectic splitting method. Higher order approximations can be obtained by composing the flows corresponding to $T(p)$ and $U(q)$ with certain coefficients obtained by solving the so-called order conditions \\cite{mclachlan02sm}. There exist in the literature a vast number of high order integrators constructed along this line (see, e.g., \\cite{blanes08sac}, \\cite{hairer06gni},  and references therein).  The non-relativistic gravitational N-body problem, in particular, belongs to this class of systems. If one considers the motion of $n+1$ particles (the Sun, with mass $m_0$, and $n$ planets with masses $m_i$, $i=1,\\ldots,n$) only affected by their mutual gravitational interaction, the corresponding equations of motion can be derived from the Hamiltonian \\begin{equation}   \\label{n-body} H = \\frac{1}{2} \\sum_{i=0}^n \\frac{\\|\\mathbf{p}_i\\|^2}{m_i} - G \\sum_{0 \\le i < j \\le n} \\frac{m_i m_j}{\\|\\mathbf{q}_i - \\mathbf{q}_j\\|}, \\end{equation} where $\\mathbf{q}_i$ and $\\mathbf{p}_i = m_i \\, \\mathbf{\\dot{q}}_i$ denote the position and momenta of the $n+1$ bodies in a barycentric reference frame.  Typically, the planets evolve around the central mass following almost Keplerian orbits, so that by an appropriate change of coordinates one can rewrite the Hamiltonian (\\ref{n-body}) as $H = H_K + H_I$, where in some sense $|H_I| \\ll |H_K|$, or equivalently, as the sum of the Keplerian motion of each planet around the central mass and a small perturbation due to the gravitational interaction between planets. Jacobi and Heliocentric coordinates constitute paradigmatic examples of canonical set of coordinates possessing this feature. Thus, the Hamiltonian (\\ref{n-body}) written as $H = H_K + H_I$ is a particular example of a near-integrable Hamiltonian system, i.e, it can be expressed as \\begin{equation}    \\label{intro.1} H(q,p;\\eps) = H^{[a]}(q,p) + \\eps H^{[b]}(q,p), \\end{equation} where $\\eps \\ll 1$ and $H^{[a]}$ is exactly integrable. It makes sense, then, to take into account this special structure when designing integration methods to approximate its dynamics. The idea consists in constructing splitting schemes as compositions of the flows corresponding to $H^{[a]}(q,p)$ and $H^{[b]}(q,p)$, assuming that they are explicitly computable or sufficiently well approximated \\cite{kinoshita91sia,wisdom91smf}. In fact, since the parameter $\\eps$ is small, it is possible to design methods which behave in practice as high order integrators with less severe restrictions concerning the order conditions than the usual split into kinetic and potential energy. This approach was systematically pursued by McLachlan \\cite{mclachlan95cmi}, obtaining families of splitting schemes of order 2 and 4 which eliminate the most relevant error terms in $\\eps$, and further analyzed by Laskar \\& Robutel \\cite{laskar01hos} in the context of planetary motion.  By incorporating the idea of processing, even more efficient schemes can be constructed for the Hamiltonian (\\ref{intro.1}) \\cite{blanes00psm}. In that case, both the kernel and the processor are taken as compositions of the flows associated with $H^{[a]}$ and $H^{[b]}$, so that the exactly symplectic character of the integration scheme is ensured. With this approach, all terms of first order in $\\eps$ in the truncation error expansion can be annihilated with the processor \\cite{mclachlan96mos,wisdom96sco}.    Although the symplectic methods developed in \\cite{laskar01hos} and \\cite{mclachlan95cmi}  for near-in\\-te\\-gra\\-ble Hamiltonian systems have proved their usefulness in long term integrations of the Solar System \\cite{laskar09eoc}, the design of new and more efficient higher order integrators is of interest for numerical simulations of its evolution over large  time spans, either by speeding up the algorithms or by providing better  accuracy in the position of the different objects. Relevant examples where the new integrators could be useful include  the numerical integration of the Solar System for more than 60 million years backward in time to cover the Palaeogene period to determine insolation quantities of the Earth and calibrate paleoclimatic data, studies of the planetary orbits over several billion years, etc. \\cite{laskar04alt}. To this purpose, it is essential that the numerical solutions obtained are not contaminated by error accumulations along the integration and that the computations are done in a reasonable time.  These long-time numerical integrations can be combined with standard techniques of classical perturbation theory, such as the expansion of  the equations of motion up to a certain order in the perturbation parameters and the use of averaging (see, e.g. \\cite{laskar86sto}).   The purpose of this work is to present new families of symplectic splitting methods specifically designed for Hamiltonian systems of the form (\\ref{intro.1}) appearing in many problems of dynamical astronomy, when one is interested in highly accurate results over a large time span. The schemes we propose will be useful in particular in the long time integration of the Solar System, both in Jacobi and Poincar\\'e Heliocentric coordinates, and are more efficient than the schemes designed in \\cite{laskar01hos} and \\cite{mclachlan95cmi}. Although they involve the computation of more elementary flows per step than other methods,  their small error terms allow to use larger steps, which results in more efficient schemes. Obtaining these new methods requires deriving previously the necessary and sufficient order conditions to be satisfied by the coefficients (which is done here in a systematic way) and then solving these polynomial equations to get the best solutions according with some appropriately chosen optimization criteria. This is discussed in more detail in sections \\ref{sec.2}, \\ref{sec.3} and the appendix, whereas in section \\ref{sec.4} we consider the application of the new schemes to the integration of the Solar System. The new methods obtained in section~\\ref{sec.3} are suitable to be applied when using Jacobi coordinates and also Poincar\\'e Heliocentric coordinates.    It is worth stressing that,  while the main motivation of this work is the long time integration of Hamiltonian problems arising in dynamical astronomy, and in particular in planetary systems, the new symplectic splitting methods obtained here can also be applied to more general perturbed differential equations arising in different fields when high accuracies are required. ",
        "conclusions": "In reference \\cite{laskar01hos}, symplectic splitting methods of generalized order $(2n,2)$ up to $n=5$ (first described in \\cite{mclachlan95cmi}) were systematically derived and tested on the Sun--Jupiter--Saturn system over 25000 years in Jacobi coordinates, observing an improvement in the accuracy with respect to the leapfrog integrator by several orders of magnitude at the same computational cost. Methods in this family have all the coefficients positive and good stability properties. Scheme $(8,2)$ in particular has been used in several long term simulations of the whole Solar System (e.g., \\cite{laskar09eoc,laskar04alt}) and corresponds to method \\texttt{ABA82} in the examples reported here. All the tests carried out in \\cite{laskar01hos} showed that the error term $\\eps^2 \\tau^3$ was the main limiting factor in the performance of the integrators, so the natural question was whether schemes of higher order (and thus already involving some negative coefficients) could be useful for integrating planetary N-body problems.  As a matter of fact, methods of order $(8,4)$ obtained in \\cite{mclachlan95cmi} do improve the performance of \\texttt{ABA82} for this problem in Jacobi coordinates, as the experiments reported in \\cite{farres12sif} show. In this work we have pursued this line of research and constructed new families of higher order splitting methods specifically oriented to the numerical integration of near-integrable Hamiltonian systems, and in particular for planetary N-body problems, both in Jacobi and Poincar\\'e Heliocentric coordinates. For this purpose, first we have derived explicitly the set of independent necessary and sufficient order conditions that splitting methods must verify to achieve a certain order of accuracy and then we have solved these equations. A non-trivial task that requires the use of homotopy continuation techniques and optimization criteria to select the most appropriate solution.  Although the new methods involve some negative coefficients, and thus one could think that their numerical stability might be compromised, they have been selected to minimize the error terms at higher orders and the sum of the absolute values of their coefficients. As a result, the size of the negative coefficients of our new methods is relatively small. In any case, the experiments reported here clearly indicate that the new methods of order $(8,6,4)$ and $(10,6,4)$ achieve accuracy up to round off error with larger step sizes than 2nd-order schemes.  There are near-integrable systems of the form (\\ref{eq:1}) where the exact flow $\\varphi_{\\tau}^{[b]}$ corresponding to the perturbation is not available. In that case, we have constructed splitting methods of the form (\\ref{eq:3approx}), where $\\widetilde{\\varphi}_{\\tau}^{[b]}$ is a 2nd-order symmetric approximation of the exact flow. This class of schemes has shown to be particularly efficient for long time integrations of N-body planetary systems in Poincar\\'e Heliocentric coordinates when the leapfrog approximation (\\ref{helio.4}) is considered.   Numerical simulations show that the efficiency of the new integrators presented here is essentially similar in both Jacobi and Heliocentric coordinates. We believe this result is worth remarking, since canonical Heliocentric variables provide very often a more convenient formulation of the problem. The improvement of the new integrators presented here (in particular, methods of order $(8,6,4)$ and $(10,6,4)$) with respect to previous schemes is most notably exhibited when they are applied for the numerical integration of the outer planets. When the whole Solar System is considered, although methods of order $(8,6,4)$ and $(10,6,4)$ still provide the best results, it is Mercury with its relatively high eccentricity and fast orbital period  which constitutes the main limiting factor in all simulations.  When designing the new methods of generalized order $(8,6,4)$ and $(10,6,4)$, we have only considered compositions with the minimum number of stages to solve all the order conditions. It might be the case, as in other contexts, that introducing more stages with additional free parameters could lead to more efficient schemes. We intend to explore this possibility and eventually collect the new methods obtained, both in the ABA and BAB classes, in our website (\\texttt{www.gicas.uji.es/software.html}).   \\subsection*{Acknowledgements}  The work of SB, FC, JM and AM has been partially supported by Ministerio de Ciencia e Innovaci\\'on (Spain) under project MTM2010-18246-C03 (co-financed by FEDER Funds of the European Union), whereas AF and JL acknowledge financial support by the FP7 GTSnext project.   \\begin{appendix}"
    },
    "1208/1208.2323_arXiv.txt": {
        "abstract": "Brown dwarfs are important objects because they may provide a missing link between stars and planets, two populations that have dramatically different formation history. In this paper, we present the candidate binaries with brown dwarf companions that are found by analyzing binary microlensing events discovered during 2004 -- 2011 observation seasons. Based on the low mass ratio criterion of $q < 0.2$, we found 7 candidate events, including OGLE-2004-BLG-035, OGLE-2004-BLG-039, OGLE-2007-BLG-006, OGLE-2007-BLG-399/MOA-2007-BLG-334, MOA-2011-BLG-104/OGLE-2011-BLG-0172, MOA-2011-BLG-149, and MOA-201-BLG-278/OGLE-2011-BLG-012N. Among them, we are able to confirm that the companions of the lenses of MOA-2011-BLG-104/OGLE-2011-BLG-0172 and MOA-2011-BLG-149 are brown dwarfs by determining the mass of the lens based on the simultaneous measurement of the Einstein radius and the lens parallax. The measured mass of the brown dwarf companions are $(0.02\\pm0.01)$ $M_{\\odot}$ and $(0.019\\pm0.002)$ $M_{\\odot}$ for MOA-2011-BLG-104/OGLE-2011-BLG-0172 and MOA-2011-BLG-149, respectively, and both companions are orbiting low mass M dwarf host stars. More microlensing brown dwarfs are expected to be detected as the number of lensing events with well covered light curves increases with new generation searches. ",
        "introduction": "Brown dwarfs are sub-stellar objects that are too low in mass to sustain hydrogen fusion reactions in their cores but much higher than planets. Studies of brown dwarfs are important because they may provide a missing link between stars and planets \\citep{kulkarni97}, two populations that have dramatically different formation history. In addition, the Galaxy may be teeming with brown dwarfs although there is no evidence for a large population of brown dwarfs from current observations \\citep{graff96, najita00, tisserand07}. Brown dwarfs had long been thought to exist based on theoretical considerations \\citep{kumar69}. However, these objects are intrinsically faint and thus the first confirmed brown dwarf was not discovered until 1995 \\citep{rebolo95, nakajima95}.  There have been 2 major methods of detecting brown dwarfs. The first is direct imaging by using ground (e.g.,\\ Schuh 2003) and space-borne (e.g.,\\ Mainzer 2011) infrared instruments. As in the case of extrasolar planets, brown dwarfs can also be indirectly discovered by detecting wobbles in the motion of the companion star or the flux decrease of the companion star caused by the brown dwarf occulting the companion star.  Microlensing is also an effective method to detect brown dwarfs. Gravitational microlensing refers to the astronomical phenomenon wherein the brightness of a star is magnified by the bending of light caused by the gravity of an intervening object (lens) between the background star (source) and an observer. Since the phenomenon occurs regardless of the lens brightness, microlensing was proposed to detect dark components of the Galaxy such as black holes, neutron stars, and brown dwarfs \\citep{paczynski86}.  Although effective, the application of the microlensing method in brown dwarf detections has been limited. The main reason for this is the difficulty in distinguishing brown dwarf events from those produced by main sequence stars. For general lensing events in which a single mass causes the brightening of a background star, the magnification of the lensed star flux depends on the projected lens-source separation, $u$, by \\begin{equation} A={ u^2+2 \\over { u(u^2+4)^{1/2} } }, \\end{equation} where the lens-source separation is normalized by the radius of the Einstein ring, $\\theta_{\\rm E}$ (Einstein radius). The lens-source separation is related to the lensing parameters of the time scale for the source to cross the Einstein radius, $t_{\\rm E}$ (Einstein time scale), the time of the closest lens-source approach, $t_0$, and the lens-source separation at the moment of the closest approach, $u_0$ (impact parameter), by \\begin{equation} u=\\left[{ \\left({t-t_0 \\over{t_{\\rm E}}}\\right)^2 + {u_0}^2 }\\right]^{1/2}. \\end{equation} Among these lensing parameters, only the Einstein time scale provides information about the physical parameters of the lens. However, it results from the combination of the lens mass, distance, and transverse speed of the relative lens-source motion and thus it is difficult to uniquely determine the mass of the lensing object. Not knowing the mass, then, it is difficult to single out brown dwarf events from those produced by stars. In principle, it is possible to measure the lens mass by additionally measuring the Einstein radius and the lens parallax, which are respectively related to the physical parameters of the lens by \\begin{equation} \\theta_{\\rm E} = \\left( \\kappa M \\pi_{\\rm rel} \\right)^{1/2} \\end{equation} and \\begin{equation} \\pi_{\\rm E} ={ \\pi_{\\rm rel} \\over{\\theta_{\\rm E} } }, \\end{equation} where $\\pi_{\\rm rel}={\\rm AU}(D_{\\rm L}^{-1}-D_{\\rm S}^{-1})$, $\\kappa=4G/(c^2{\\rm AU})$, $M$ is the mass of the lens, and $D_{\\rm L}$ and $D_{\\rm S}$ are the distances to the lens and source star, respectively. The Einstein radius is measured by analyzing the distortion of the lensing light curve affected by the finite size of the source star \\citep{nemiroff94, witt94, gould94}, while the lens parallax is measured by analyzing the deviation in a lensing light curve caused by the deviation of the relative lens-source motion from a rectilinear one due to the change of the observer's position induced by the orbital motion of the Earth around the Sun \\citep{gould92}. Unfortunately, the chance of simultaneously measuring $\\theta_{\\rm E}$ and $\\pi_{\\rm E}$ is very low for a single-lens event. The finite source effect is important only for very high magnification events, for which the lens-source separation is comparable to the source size \\citep{choi12}, but these events are very rare. The parallax effect is important only for events with time scales that are a significant fraction of the orbital period of the Earth, i.e.\\ 1 yr. However, such long time-scale events are unlikely to be produced by low mass brown dwarfs because $t_{\\rm E}\\propto M^{1/2}$. As a result, there exist only two cases for which the lens of a single lensing event was identified as a brown dwarf \\citep{smith03, gould09} \\footnote{Among these two cases, the microlensing brown dwarf discovered by \\citet{gould09} is exception that provides the rule because the lens mass could be determined not by measuring the lens parallax induced by the Earth's orbital motion but from the ``terrestrial parallax''.}.      In binary lensing events, in which the lens is composed of two masses, on the other hand, the chance to identify the lens as a brown dwarf is relatively high. There are several reasons for this. First, for binary lensing events, it is possible to routinely measure the mass ratio, $q$, of the lens components. Then, brown dwarf candidates can be sorted out based on the first order criterion of small mass ratios. Second, the majority of binary lensing events involves caustic crossings \\citep{schneider86, mao91}. The caustics represent the positions on the source plane at which the lensing magnification of a point source becomes infinite. The caustic-crossing part of a lensing light curve varies depending on the source size. Analysis of the light curve affected by this finite source effect allows one to measure the Einstein radius, which allows for better constraint on the lens mass. Third, the chance to measure the lens parallax is high, too. This is because the duration of an event corresponds to the total mass of the binary lens system not the low mass component and thus the average time scale of a binary lens event is longer than that of a single lens event.   Despite these advantages, there exists only one event (OGLE-2008-BLG-510) for which the companion of a binary lens is known to be a brown dwarf candidate \\citep{bozza12}. The main reason for the rarity of brown dwarf events is the difficulty in analyzing binary lensing light curves due to the complexity of the light curves. For binary lensing events in which the mass ratio between the lens components is very low as in the case of a binary composed of a star and a planet, the signature of the low mass companion appears as a short term deviation on the top of an otherwise smooth and symmetric single lensing light curve of the primary. As a result, identification and analysis of planetary events are relatively simple, resulting in more than 20 identified microlensing planets. However, the perturbative nature of the companion breaks down and the light curve becomes very complex for binaries with mass ratios greater than $q\\sim0.1$, which corresponds to the mass ratio of a typical binary lens composed of a brown dwarf and a star. This introduces difficulties in the immediate identification of a brown dwarf event. A firm identification of a brown dwarf event requires detailed modeling of lensing light curves. Unfortunately, modeling light curves of binary lensing events is a difficult task due to the complexity of $\\chi^2$ surface in the parameter space combined with the heavy computation required for analysis. As a result, detailed modeling of binary events was conducted for a small subset of events, resulting in a small number of identified brown dwarfs. In recent years, however, there has been great progress in the analysis of binary events with the development of analysis methods based on advanced logic combined with improved computing power. This progress made the analysis not only more precise but also faster. As a result, nearly all binary lensing events are routinely modeled in real time with the progress of events in current microlensing experiments \\citep{ryu10, bozza12}.  In this paper, we present brown dwarf candidates found from the analyses of binary lensing events detected from microlensing experiments conducted from 2004 to 2011. In \\S 2, we describe the criteria of choosing brown dwarf candidates, observation of events, and data reduction. In \\S 3, we describe the procedure of modeling. In \\S 4, we present the results of the analysis. In \\S 5, we draw our conclusions.    \\begin{deluxetable*}{lrrrr} \\tablecaption{Coordinates of Events\\label{table:one}} \\tablewidth{0pt} \\tablehead{ \\multicolumn{1}{c}{event} & \\multicolumn{1}{c}{RA (J2000)} & \\multicolumn{1}{c}{DEC (J2000)} & \\multicolumn{1}{c}{$l$} & \\multicolumn{1}{c}{$b$} } \\startdata OGLE-2004-BLG-035                   & 17$^{\\rm h}$48$^{\\rm m}$43$^{\\rm s}$\\hskip-2pt.16 & -35$^\\circ$57${\\rm '}$45${\\rm ''}$\\hskip-2pt.9 & 354.32$^\\circ$ & -4.19$^\\circ$ \\\\ OGLE-2004-BLG-039                   & 17$^{\\rm h}$53$^{\\rm m}$47$^{\\rm s}$\\hskip-2pt.58 & -30$^\\circ$52${\\rm '}$11${\\rm ''}$\\hskip-2pt.7 & 359.25$^\\circ$ & -2.51$^\\circ$ \\\\ OGLE-2007-BLG-006                   & 18$^{\\rm h}$02$^{\\rm m}$52$^{\\rm s}$\\hskip-2pt.50 & -29$^\\circ$15${\\rm '}$11${\\rm ''}$\\hskip-2pt.8 &   1.63$^\\circ$ & -3.41$^\\circ$ \\\\ OGLE-2007-BLG-399/MOA-2007-BLG-334  & 17$^{\\rm h}$45$^{\\rm m}$35$^{\\rm s}$\\hskip-2pt.79 & -34$^\\circ$54${\\rm '}$37${\\rm ''}$\\hskip-2pt.5 & 354.89$^\\circ$ & -3.10$^\\circ$ \\\\ MOA-2011-BLG-104/OGLE-2011-BLG-0172 & 17$^{\\rm h}$54$^{\\rm m}$22$^{\\rm s}$\\hskip-2pt.48 & -29$^\\circ$50${\\rm '}$01${\\rm ''}$\\hskip-2pt.7 &   0.21$^\\circ$ & -2.10$^\\circ$ \\\\ MOA-2011-BLG-149                    & 17$^{\\rm h}$56$^{\\rm m}$47$^{\\rm s}$\\hskip-2pt.69 & -31$^\\circ$16${\\rm '}$04${\\rm ''}$\\hskip-2pt.7 & 359.23$^\\circ$ & -3.27$^\\circ$ \\\\ MOA-2011-BLG-278/OGLE-2011-BLG-012N & 17$^{\\rm h}$54$^{\\rm m}$11$^{\\rm s}$\\hskip-2pt.32 & -30$^\\circ$05${\\rm '}$21${\\rm ''}$\\hskip-2pt.6 & 359.97$^\\circ$ & -2.19$^\\circ$ \\enddata \\end{deluxetable*}     \\begin{deluxetable}{ll} \\tablecaption{Telescopes\\label{table:two}} \\tablewidth{0pt} \\tablehead{ \\multicolumn{1}{c}{event} & \\multicolumn{1}{c}{telescopes} } \\startdata OGLE-2004-BLG-035   & OGLE, 1.3 m Warsaw, Las Campanas, Chile  \\\\ \\hline OGLE-2004-BLG-039   & OGLE, 1.3 m Warsaw, Las Campanas, Chile  \\\\ \\hline OGLE-2007-BLG-006   & OGLE, 1.3 m Warsaw, Las Campanas, Chile  \\\\ & $\\mu$FUN, 1.3 m SMARTS, CTIO, Chile      \\\\ & $\\mu$FUN, 0.4 m Auckland, New Zealand    \\\\ & $\\mu$FUN, 0.4 m FCO, New Zealand         \\\\ \\hline OGLE-2007-BLG-399   & OGLE, 1.3 m Warsaw, Las Campanas, Chile  \\\\ /MOA-2007-BLG-334   & MOA, 2.0 m Mt. John, New Zealand         \\\\ & $\\mu$FUN, 1.3 m SMARTS, CTIO, Chile      \\\\ & PLANET, 1.0 m SAAO, South Africa         \\\\ \\hline MOA-2011-BLG-104    & OGLE, 1.3 m Warsaw, Las Campanas, Chile  \\\\ /OGLE-2011-BLG-0172 & MOA, 2.0 m Mt. John, New Zealand         \\\\ & $\\mu$FUN, 1.3 m SMARTS, CTIO, Chile      \\\\ & $\\mu$FUN, 0.4 m Auckland, New Zealand    \\\\ & $\\mu$FUN, 0.4 m FCO, New Zealand         \\\\ & $\\mu$FUN, 0.5 m Wise, Israel             \\\\ & $\\mu$FUN, 0.3 m PEST, Australia         \\\\ & RoboNet,  2.0 m FTN, Hawaii              \\\\ & RoboNet,  2.0 m FTS, Australia           \\\\ \\hline MOA-2011-BLG-149    & MOA, 2.0 m Mt. John, New Zealand         \\\\ & OGLE, 1.3 m Warsaw, Las Campanas, Chile  \\\\ & $\\mu$FUN, 1.3 m SMARTS, CTIO, Chile      \\\\ & $\\mu$FUN, 0.4 m Auckland, New Zealand    \\\\ & RoboNet, 2.0 m FTS, Australia            \\\\ & MiNDSTEp, 1.54 m Danish, La Silla, Chile \\\\ \\hline MOA-2011-BLG-278    & MOA, 2.0 m Mt. John, New Zealand         \\\\ /OGLE-2011-BLG-012N & OGLE, 1.3 m Warsaw, Las Campanas, Chile  \\\\ & $\\mu$FUN, 1.3 m SMARTS, CTIO, Chile      \\\\ & $\\mu$FUN, 0.4 m Auckland, New Zealand    \\\\ & $\\mu$FUN, 0.4 m FCO, New Zealand         \\\\ & $\\mu$FUN, 0.4 m Kumeu, New Zealand       \\\\ & $\\mu$FUN, 1.0 m Lemmon, Arizona          \\\\ & RoboNet, 2.0 m FTS, Australia            \\\\ & MiNDSTEp, 1.54 m Danish, La Silla, Chile \\enddata \\tablecomments{ CTIO: Cerro Tololo Inter-American Observatory; FCO: Farm Cove Observatory; SAAO: South Africa Astronomical Observatory; PEST: Perth Extrasolar Survey Telescope; FTN: Faulkes North; FTS: Faulkes South } \\end{deluxetable} ",
        "conclusions": "We searched for candidate binaries with brown dwarf companions by analyzing binary microlensing events discovered during the 2004 -- 2011 observing seasons. Under the low mass ratio criterion of $q < 0.2$, we found 7 candidate events, including OGLE-2004-BLG-035, OGLE-2004-BLG-039, OGLE-2007-BLG-006, OGLE-2007-BLG-399/MOA-2007-BLG-334, MOA-2011-BLG-104/OGLE-2011-BLG-0172, MOA-2011-BLG-149, and MOA-201-BLG-278/OGLE-2011-BLG-012N. Among them, we confirmed that the companions of MOA-2011-BLG-104/OGLE-2011-BLG-0172 and MOA-2011-BLG-149 were brown dwarfs by measuring the lens masses.  The number of microlensing brown dwarfs is expected to increase with the improvement of lensing surveys and analysis technique. With the adoption of a new wide field camera, the number of lensing events detected by the OGLE survey is increased by a factor more than 2. In addition, a new survey based on a network of 3 telescopes located at three different locations of the Southern Hemisphere (Korea Microlensing Telescope Network) is planned to operate from the 2014 season. Along with the improvement in the observational side, the analysis technique has greatly improved for the last several years, enabling prompt and precise analysis of a large number of anomalous events. With this improvement, the number of microlensing brown dwarfs will rapidly increase, making microlensing one of the major methods in the discovery of brown dwarfs.    \\mbox{}"
    },
    "1208/1208.2862_arXiv.txt": {
        "abstract": "We have analysed the combined Greenwich and Solar Optical Observing Network (SOON) sunspot group data during the period of 1874\\,--\\,2011 and determined variations in the annual numbers (counts) of the small (maximum area $A_{\\rm M} < 100$ millionth of solar hemisphere, msh), large ($100 \\le A_{\\rm M} < 300$ msh), and big ($A_{\\rm M} \\ge 300$ msh) spot groups. We  found that the amplitude of an even-numbered   cycle of  the number of large groups is smaller  than that of  its immediately following odd-numbered cycle. This is consistent with  the well known Gnevyshev and Ohl rule or G-O rule of solar cycles, generally described  by using the Z\\\"urich sunspot number ($R_{\\rm Z}$).  During cycles 12\\,--\\,21 the G-O rule holds good for the variation in the number of small groups  also, but it is violated by cycle pair (22, 23) as in the case of $R_{\\rm Z}$.   This behavior of the variations in the small groups is largely  responsible  for the anomalous behavior of  $R_{\\rm Z}$ in cycle pair (22, 23). It is also found that   the amplitude of an  odd-numbered cycle of the number of small groups is larger   than that of  its immediately following even-numbered cycle. This can be called as `reverse G-O rule'. In the case of the number of the big groups, both   cycle pairs (12, 13) and (22, 23)   violated the G-O rule. In many cycles the positions of the peaks of the small, large, and big groups   are different and considerably differ with respect to the corresponding positions of the $R_{\\rm Z}$ peaks. In the case  of cycle~23, the corresponding cycles of the small and large   groups  are largely symmetric/less asymmetric (Waldmeier effect is weak/absent) with their maxima taking place two years later than that of $R_{\\rm Z}$. The corresponding cycle of the big groups is more asymmetric (strong Waldmeier effect) with its maximum epoch taking place at the  same time as  that of $R_{\\rm Z}$. ",
        "introduction": "Studies on variations in solar activity are important for understanding the mechanism behind the solar activity and solar cycle,  and also for predicting the level of activity~\\cite{hat09,pet10}. The properties of solar cycle are generally described by the Z\\\"urich or international sunspot number, $R_{\\rm Z} =  k (10 g + f)$, where $k$ is a correction factor for the observer, $g$ is the number of identified sunspot groups, and $f$ is the number of individual sunspots. Several other solar activity indices  well correlate with $R_{\\rm Z}$~\\cite{hw04}. However, there are noticeable differences in the epochs of the peaks of $R_{\\rm Z}$ and  other activity indices  in  some solar cycles ($e.g.$, see~\\inlinecite{rr08} and references therein). It seems there are also considerable differences between the  epochs of the maxima of sunspot cycles and the corresponding  cyclic variations in the sunspot field strength ~\\cite{pevt11}. It is believed that the area of a sunspot or a sunspot group  has a better physical significance   than $R_{\\rm Z}$ because the area is a better measure (proxy) of  solar magnetic flux than $R_{\\rm Z}$~\\cite{dg06}; an area of 130 msh (millionths of solar hemisphere; 1 msh $\\approx 3\\times 10^6$ km$^2$) corresponds approximately to $10^{22}$ Mx (maxwell)~\\cite{ws89}. There exists a high correlation between  solar cycle variations of $R_{\\rm Z}$ and sunspot area~\\cite{hwr02,hw04}. There are some minor but noticeable  differences in the variations of  $R_{\\rm Z}$  and sunspot group area. \\inlinecite{dgt08} reported that the well-known Waldmeier effect  (inverse relationship between the rise time and the amplitude of a cycle) does not exist in the case of sunspot area. The  epochs of the maxima of some cycles of sunspot number and sunspot area are  different. For example, in case of  cycle~23, the epoch of maximum of the sunspot number was in 2000, whereas the maximum of the sunspot area  was in 2002~\\cite{kbr10,jj12}.   The sunspot cycles are numbered in chronological order from the cycle that started from the year 1755, and are  known as Waldmeier cycle numbers. Many  characteristics of sunspot cycles are known (see \\opencite{hw04}). Usually  the  time series  $R_{\\rm Z}$ is used to reveal most of the characteristics of the solar cycle. However, Hoyt and Schatten (\\citeyear{hk88a}, \\citeyear{hk88b}) devised a number index based solely on the number of observed sunspot groups. The group sunspot number, $R_{\\rm G}$, gives a more complete and longer data set than  $R_{\\rm Z}$~\\cite{hwr02,hw04}, but the time series  $R_{\\rm G}$ ended in 1995. None of the characteristics of the solar cycle is fully understood so far. One of the most  prominent and  fundamental characteristics is the differences in the amplitudes and the lengths of different solar cycles. According to the well-known Gnevyshev-Ohl rule~\\cite{gh48} or G-O rule, the amplitude of an odd-numbered cycle is higher than that of the preceding even-numbered cycle. However, there are instances of violation of this rule, {\\it viz}, cycle pairs (4, 5), (8, 9), and (22, 23). So far no plausible method is available to predict the violation of the G-O rule (except that it may be possible from the epochs of the retrograde motion of the Sun about the solar system barycenter, as suggested by ~\\inlinecite{jj05}). Using the data on sunspot groups during the period 1879\\,--\\,2004, \\inlinecite{jbu05a} found that the solar equatorial rotation rate during an odd--numbered sunspot cycle   correlates well with  the equatorial rotation rate of the preceding even--numbered sunspot cycle, which is similar to the G-O rule in sunspot activity. They also found that  the latitudinal gradient of  solar rotation during  an even--numbered cycle correlates well with   that of the preceding odd--numbered cycles. These results seem to imply that the G-O rule is related to   the basic mechanism of solar activity and solar cycle.   Typical sizes of sunspots range from 10 to $10^3$ msh. Although single sunspots are common,  the majority of sunspots belong to  sunspot groups. Sunspot groups are often large and complex. It is generally believed that large sunspot groups also  live long. In fact, there is a rule of  proportionality between the maximum area ($A_{\\rm M}$) of a sunspot group and its life time ($T$) (first noticed by~\\inlinecite{gne38} and formulated by~\\inlinecite{wald55}; see also~\\opencite{pvd97}): $\\frac {A_{\\rm M}} {T} \\approx   10\\ {\\rm msh\\ day}^{-1}$. However, the relationship between the area and the life time of sunspot groups may be exponential rather than linear~\\cite{jj03a}. Properties such as the  rotation rate, meridional motion, tilt angle, {\\it etc.}, of  sunspots and sunspot groups depend on their life time and size as well as their age ($e.g.$,~\\opencite{war65}, \\citeyear{war66}; \\opencite{hgg84}; \\opencite{bal86}; \\opencite{jg97}; \\opencite{jj99}; \\opencite{siva07}). Studies on these properties of sunspot groups  may provide information on the subsurface dynamics of the Sun ($e.g.$, \\opencite{how96}; \\opencite{jg97}; \\opencite{kmh02}; \\opencite{siva03}, 2007, 2010). Therefore, the studies on the variations in the numbers (counts)  of sunspots and sunspot groups in different sizes look to be important for understanding the basic mechanism of solar activity and solar cycle, and also the relationship between   sunspots and other activity indices ($e.g.$, \\opencite{kilc11}; \\opencite{cl12}). In the present paper we have analyzed the sunspot group data during the period 1874\\,--\\,2011  and studied the variations in the annual numbers of sunspot groups of different sizes. Particularly  we have concentrated on the G-O rule and the Waldmeier effect in the variations of small and large sunspot groups, and their implications.  In the next section we will describe the data and the method of analysis. In Section~3 we will describe the results and in Section~4 we will  present conclusions and a brief discussion. ",
        "conclusions": "From the above analyses of  a large data set on sunspot groups the following conclusions can be drawn: \\begin{enumerate} \\item  The amplitude of an even--numbered   cycle of  NLGs is smaller  than that of  its immediately following odd--numbered cycle. This is consistent with  the well-known G-O rule of solar cycles. Obviously,  the amplitude of cycle~22 of NLGs is smaller than that of cycle~23. This is in  line with the conclusion drawn by~\\inlinecite{kilc11}. \\item NSGs  also satisfies the even--odd cycle rule, but cycle pair (22, 23) violated the G-O rule, {\\it i.e.}, in this cycle pair the behavior of NSGs is  similar to that of   $R_{\\rm Z}$. It is also found that   the amplitude of an  odd--numbered cycle of NSGs  is larger   than that of  its immediately following even--numbered cycle. This can be called as a `reverse G-O rule'. \\item In the case of NBGs,   cycle pairs (12, 13) and (22, 23) show similar behavior, {\\it i.e.}, both violated the even-odd cycle rule. The violation of the G-O rule in NBGs   is not reflected in $R_{\\rm Z}$  because contributions from   NBGs to $R_{\\rm Z}$ are relatively small. \\item In many cycles the positions of the peaks of NSGs, NLGs, and NBGs  are different, and they also deviate considerably from the corresponding peak positions of  $R_{\\rm Z}$. In the case of cycle~23, the maxima of NSGs and NLGs  are at 2002, whereas the maximum of NBGs is at 2000, $i.e.$, at the  same epoch of the maximum of $R_{\\rm Z}$. The corresponding  cycles in NSGs and NLGs  are largely symmetric or less asymmetric (the Waldmeier effect is weak or absent), and  the corresponding cycle in NBGs  is more asymmetric (strong Waldmeier effect). \\item The amplitude of the long-term variation is large in the case of NSGs,  and its  period  ($\\approx$ 180 years) is also long in this case. An approximate 90-year cycle is seen in NBGs. The long-term trend in the number of small groups  implies that the current cycle~24 is weak. \\end{enumerate}   The studies of rotation rates of sunspot  groups of different life times and sizes indicated that the magnetic structures of small and large/big  groups anchor near the surface (near 0.95$R_\\odot$, here $R_\\odot$ is the radius of the Sun) and relatively deeper layers  (even reach near to the base of the convection zone) of the Sun's convection zone, respectively~(\\opencite{jg97}; \\opencite{kmh02}; \\opencite{siva03}, 2004). Such a study also suggested that small sunspot groups may be the fragmented or the branched parts of the large/big sunspot groups~\\cite{jj03a}. In other words,  large/big sunspot groups may be the products of a deep dynamo mechanism, whereas  small sunspot groups may be the product of a surface dynamo mechanism. Thus,  our conclusions 1 and 2 above  imply that the G-O rule represents a deep rooted global property of the solar cycle. Because of their large numbers the small sunspot groups mainly contributed to the behavior of cycle pair (22, 23) to violate the G-O rule. That is, the violation of the G-O rule in $R_{\\rm Z}$ may be largely related to the surface (local) dynamo,  for  which convection and  meridional flows may  have major roles (cancellation of magnetic flux). The `reverse G-O rule' found in the variation of the number of small groups (conclusion 2 above) is also related to the surface (local) dynamo mechanism.  The existence of approximate 90-year and 180-year periodicities  has been known in sunspot activity. In the present analysis they are seen (visualized) in the variations of the numbers of large/big sunspot groups and of small sunspot groups, respectively (conclusion 5 above), suggesting that the former and latter are  related to the dynamics of the deeper  and the near-surface layers, respectively.  A `reverse G-O rule'~\\cite{jbu05a}  and  a 90-year periodicity~\\cite{jj03b,jbu05b} were also found in the latitude gradient of solar rotation determined from the sunspot group data. Conclusions~2 and 5 above indicate that  rotations of  small and  big  sunspot groups largely contribute to  the reverse G-O rule and the 90-year periodicity of the latitude gradient, respectively. Conclusion~2 above also indicates  that in terms of  the number of SSGs the amplitude of the current cycle~24 will be  smaller  than that of the previous cycle~23.  The maximum in $R_{\\rm Z}$ of  solar cycle is not smooth sometimes. Two or more peaks can be identified during the solar maxima and are called Gnevyshev peaks, because this splitting of activity was identified for the first time by Gnevyshev~(\\citeyear{gne67}, \\citeyear{gne77}). The time interval between these peaks, where the level of activity is relatively low, is known as the Gnevyshev gap (see the review by~\\opencite{sto03}). \\inlinecite{gopal03} found that during cycle~23 the  occurrence rate of coronal mass ejections (CMEs) peaked two years after the maximum epoch (2000) of this cycle. \\inlinecite{det04} reported that not only $R_{\\rm Z}$ but also all other  activity  indices, including facular area, F10.7 flux, {\\it etc.}, show a double peak structure near the maximum of solar cycle 23, and these  have the highest peak in the year 2002. \\inlinecite{kbr10} showed that the occurrence peak of CMEs is close to the peak of the sunspot  group area.  More powerful flares seem   to occur after the maximum epoch of $R_{\\rm Z}$ and the maximum of the annual number of X-class flares took  place  close to 2002~\\cite{tan11}. \\inlinecite{jj12} found that the amounts (annual rates)  of growth and decay of  magnetic flux  in sunspot groups  in a given time interval (year) correlate well with  the amount of magnetic flux available in that interval. Hence, they have claimed that the solar cycle variation in the decay of sunspot groups has a substantial contribution to the solar cycle variations in the solar energetic phenomena and the total solar irradiance (TSI). \\inlinecite{kilc11} suggested  that the excess of large sunspot groups  during the declining phase of cycle~23 is responsible for the occurrence rate of CMEs and other activity indices to reach their maxima in the year 2002. Our results  (conclusion 2 above) indicate that there is  a considerable contribution from the small sunspot groups concerning the two--year delay in the maximum of CME occurrence rate and other activity indices/energetic phenomena. That is, as suggested in   \\inlinecite{jj12}, TSI and   solar energetic phenomena such as flares and CMEs seem to be largely related to the more evolved flux of sunspot groups."
    },
    "1208/1208.0797_arXiv.txt": {
        "abstract": "We construct models of universe with a generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n})c^2$ having a linear component and a polytropic component. The linear equation of state $p=\\alpha\\rho c^2$ describes radiation ($\\alpha=1/3$), pressureless matter ($\\alpha=0$), stiff matter ($\\alpha=1$), and vacuum energy ($\\alpha=-1$). The polytropic equation of state $p=k\\rho^{1+1/n} c^2$ may be due to Bose-Einstein condensates with repulsive ($k>0$) or attractive ($k<0$) self-interaction, or have another origin. In this paper, we consider positive indices $n>0$. In that case, the polytropic component dominates in the early universe where the density is high. For $\\alpha=1/3$, $n=1$ and $k=-4/(3\\rho_P)$, we obtain a model of early universe describing the transition from a pre-radiation era to the radiation era.  The universe exists at any time in the past and there is no singularity. However, for $t<0$, its size is less than the Planck length $l_P=1.62\\, 10^{-35}\\, {\\rm m}$. In this model, the universe undergoes an inflationary expansion with the Planck density $\\rho_P=5.16 \\, 10^{99}\\, {\\rm g}/{\\rm m}^3$ that brings it to a size $a_1=2.61\\, 10^{-6}\\, {\\rm m}$ at $t_1=1.25\\, 10^{-42}\\, {\\rm s}$ (about $20$ Planck times $t_P$). For $\\alpha=1/3$, $n=1$ and $k=4/(3\\rho_P)$, we obtain a model of early universe with a new form of primordial singularity: The universe starts at $t=0$ with an infinite density and a finite radius $a=a_1$. Actually, this universe becomes physical at a time $t_i=8.32\\, 10^{-45}\\, {\\rm s}$ from which the velocity of sound is less than the speed of light. When $a\\gg a_1$, the universe evolves like in the standard model. We describe the transition from the pre-radiation era to the radiation era by analogy with a second order phase transition where the Planck constant $\\hbar$ plays the role of finite size effects (the standard Big Bang theory is recovered for $\\hbar=0$). ",
        "introduction": "The discovery of the expansion of the universe and the concept of a singularity from which the universe emerged (Big Bang) has an interesting history \\cite{luminet,nb}. In this introduction, we provide a brief review of the early development of modern cosmology, describing  very elementary (yet fundamental) cosmological models that will play an important role in our study.    In 1917, Einstein \\cite{einsteincosmo} applied the equations of general relativity to cosmology, assuming that the universe is homogeneous and isotropic. On the ground of cultural and philosophical beliefs, he considered a static universe. In order to accommodate for such a solution, he had to introduce a cosmological constant $\\Lambda$ in the equations of general relativity\\footnote{As a preamble to his paper, Einstein \\cite{einsteincosmo} considered the Newtonian approximation and replaced the Poisson equation by an equation of the form $\\Delta\\Phi-(\\Lambda/c^2)\\Phi=4\\pi G\\rho$, so as to obtain a static solution $\\Phi=-4\\pi G\\rho c^2/\\Lambda$. In Einstein's belief, the Newtonian effect of the cosmological constant was to shield the gravitational interaction on a distance $c/\\sqrt{\\Lambda}$. This is actually  incorrect. Lema\\^itre \\cite{lemaitrecosmo} was the first to understand that the cosmological constant can be interpreted as a force of ``cosmic repulsion''. In the Newtonian context, the  modified Poisson equation including the cosmological constant is $\\Delta\\Phi=4\\pi G\\rho-\\Lambda$, leading to the static solution $\\rho=\\Lambda/4\\pi G$.}. He obtained a model of the universe with a positive curvature $k=+1$ that is finite (yet unbounded) with a ``radius'' $a_E=c/\\sqrt{\\Lambda}$, a density $\\rho_E=\\Lambda/4\\pi G$, and a mass $M_E=\\pi c^3/2G\\sqrt{\\Lambda}$.  The same year, de Sitter \\cite{deSitter1,deSitter2}  discovered  another static solution of the Einstein equations (according to the review of Robertson \\cite{robertson}, this solution was suggested to de Sitter by Ehrenfest). This universe is empty ($\\rho=0$) and has a radius $a_{dS}=c\\sqrt{3/\\Lambda}$.  In this model, the stars and nebulae are just ``test particles'' that do not curve the universe.  De Sitter managed to predict that the light emitted  by these sources should be redshifted because (i) they move away from the observer due to the effect of the cosmological constant, and (ii) there is a slowing down of time with increasing distance from the observer (de Sitter effect).  He mentioned that some observations (presumably those made by Slipher since 1912) seem to confirm that claim. He also predicted a linear, or quadratic, relationship between the redshift and the distance. However, his interpretation of the redshifts remained enigmatic for a long time. Therefore, the early cosmological debate revolved around two deficient models: The model of Einstein which was unable to explain redshifts, and the model of de Sitter which contained no matter and was rather mysterious. Both were built on the assumption that the universe is static.  In 1922 and 1924, Friedmann \\cite{friedmann1,friedmann2} discovered non-static cosmological solutions of the Einstein equations describing the temporal variation of space. He considered the case of spherical (positive  curvature $k=+1$) and hyperbolic (negative curvature $k=-1$) universes. He described expanding and contracting solutions, as well as cyclic solutions, and foresaw the possible beginning of the universe in a singularity. He obtained for the first time an estimate of the age of the universe of the order of $10$ billion years. At that time, Einstein did not believe in the significance of these non stationary solutions. He first argued that the Friedmann solutions were not compatible with the field equations of general relativity \\cite{einsteincritique1}, then retracted his statement \\cite{einsteincritique2} but remained convinced that the Friedmann solutions are not physically relevant \\cite{luminet,nb}. Apart from Einstein, who criticized them, the papers of Friedmann remained unnoticed by the physical community until 1930.  In 1922, Lanczos \\cite{lanczos1922} found a non-stationary form of de Sitter's solution in which the scale factor increases exponentially rapidly. He had the key to an expanding universe but did not recognize its physical significance.  In 1923, Weyl \\cite{weyl} published a paper in which he tried to find his own way to de Sitter's redshift. He noted that the particles in de Sitter's universe move away from each other with a velocity proportional to their distance, but did not provide an explicit system of coordinates to make his claim more precise. In 1925, Lema\\^itre \\cite{lemaitre1925} found a system of coordinates in which de Sitter's model is non-static and leads to a linear relationship between the velocity (redshift) and the distance. This model explains the redshifts of distant objects in an unambiguous way. This model is Euclidean $k=0$ (which was problematic for Lema\\^itre) and the scale factor increases exponentially rapidly with time due to the cosmological constant.     In 1927,  Lema\\^itre \\cite{lemaitre1927} rediscovered independently from Friedmann the non-stationary cosmological solutions of the Einstein equations. In addition, having spent two years in Harvard and at the MIT, he was aware of cosmological observations and related the expansion of space predicted by the theory of general relativity to the recession motion of galaxies observed by Slipher, Hubble and  Str\\\"omberg. He interpreted the redshift of the nebulae as a consequence of the expansion of the universe, predicted a linear relationship between the recession velocity and the distance (the nowdays called ``Hubble relationship'') and, using available data \\cite{stromberg}, obtained an estimate of the ``Hubble constant'' $H=\\dot a/a=v/r$ two years before Hubble \\cite{hubble}\\footnote{As pointed out by Luminet \\cite{luminet} and Nussbaumer \\& Bieri \\cite{nb}, the Hubble law and the discovery of the expansion of the universe should be attributed in part to Lema\\^itre, especially if we consider that Hubble did not relate his observations to the expansion of space in the theory of general relativity (and was later reluctant in accepting this idea), while this was Lema\\^itre's main motivation. Indeed, Hubble interpreted the redshifts in terms of de Sitter's effect in a static universe, while Lema\\^itre was the first to connect them to the expansion of the universe. Unfortunately, the paper of Lema\\^itre was totally ignored at that time because it was published in French in a rather inaccessible Belgian journal. His paper was translated in English in 1931 \\cite{lemaitre1931}, but for unknown reasons, the important paragraph of the paper where Lema\\^itre gives the relation of proportionality between the recession velocity and the distance, and extracts the Hubble constant, is condensed into a single sentence. Because of this omission, Lema\\^itre is not recognized on the same footing as Hubble for being the discoverer of the expansion of the universe \\cite{luminet,nb}.}.  The work of Lema\\^itre, like the one of his predecessor Friedmann, was still not appreciated by Einstein who found it ``abominable'' \\cite{luminet,nb}.  In  1928 and 1929, Robertson \\cite{robertson1928,robertson1929} reinterpreted de Sitter's solution dynamically. However, like Lanczos \\cite{lanczos1922} and Weyl \\cite{weyl}, he did not ``see'' the expanding universe.  In 1930, Eddington, who was working with McVittie on possible dynamical solutions of the cosmological Einstein equations (following a discussion that he had with de Sitter during a meeting of the Royal Astronomical Society \\cite{deSitter1930a}), was alerted by Lema\\^itre about his own results.  He publicized them in a paper \\cite{eddington} in which he showed that the Einstein static universe is unstable (a result that was implicit in the paper of Lema\\^itre). He also emphasized a  non-singular solution of Lema\\^itre in which the universe starts from the static solution of Einstein for $t\\rightarrow -\\infty$, then expands (this model is often called the Eddington-Lema\\^itre model). In parallel, de Sitter advertised Lema\\^itre's discovery of an expanding universe in several publications \\cite{sitter1930b,sitter1930c}. Together with the supporting observations of Hubble \\cite{hubble}, and independent theoretical work by Robertson \\cite{robertson1928,robertson1929}, this really meant the demise of the static universe.  In 1931, Einstein was forced to admit the theoretical and observational reality of the expansion of the universe. He took an extreme position and renounced to the cosmological constant \\cite{einsteinzero} that he considered  being  the ``biggest blunder'' of his life (in contrast, Lema\\^itre \\cite{lemaitrecosmo} always defended the cosmological constant as one of the most fundamental ingredients of modern cosmology).  In 1932, Einstein published with  de Sitter a one-page paper \\cite{eds}  in which they consider an  expanding universe in a Euclidean space (zero curvature $k=0$) without pressure ($p=0$) and without cosmological constant ($\\Lambda=0$). This is the celebrated Einstein-de Sitter (EdS) solution that describes the matter era. The main motivation of these authors was to show that an expansion of the universe is possible even without the cosmological constant.   In the meantime, Lema\\^itre had gone one step further. If the universe is expanding, it may have started from a {\\it singularity} in the past. In 1931, he proposed a model \\cite{lemaitresingular} in which the universe begins its expansion from a singular initial state\\footnote{Eddington \\cite{eddington} also found singular solutions to the cosmological equations but rejected them as being unphysical.  In 1931, he wrote: ``Philosophically, the notion of a beginning of the present order of Nature is repugnant to me'' \\cite{eddington1931}.}, or more precisely, from a near-singularity that  he called the ``primeval atom'' \\cite{lemaitreNewton}\\footnote{Lema\\^itre \\cite{lemaitre1933} also considered a model of universe in perpetual oscillation, undergoing cycles of expansion and contraction. He poetically  called it ``phoenix-universe''. However, he finally rejected this model at regret because he found that the duration of a cycle is too short. Cyclic universes were also discussed by Friedmann \\cite{friedmann1,friedmann2} and Einstein \\cite{einsteinzero}.}. This expansion is followed by a phase of stagnation during which the universe remains close to the static Einstein solution. In Lema\\^itre's model, it is during this phase of stagnation that galaxies form. At later times, the universe enters in the de Sitter phase of exponential expansion. A new controversy started with Einstein who did not believe in the singular model of Lema\\^itre that implies that the universe has a ``beginning''. This controversy remained until the end of their life. On the other hand, in 1960, Fred Hoyle, who promoted  a ``steady state theory'' \\cite{hoyle} in which the universe has always been identical to itself and in which matter is created spontaneously, and continuously, made fun of Lema\\^itre during a meeting in Pasadena by introducing him as the ``Big Bang man'' \\cite{luminet}. This is how the term ``Big Bang'' was introduced in cosmology and became popular.  Up to the mid-1960's, it was not clear whether the early universe had been hot or cold.  The discovery  by Penzias and Wilson \\cite{pw} of the $2.7\\, {\\rm K}$ cosmic microwave background radiation (CMB), whose existence had been predicted by the hot universe theory \\cite{gamow,dn,dicke}, immediately led widespread acceptance to the Big Bang theory. According to that theory, the early universe was filled with an ultrarelativistic gas of photons, electrons, positrons, quarks, antiquarks, etc. This is described by an equation of state $p=\\rho c^2/3$ from which it is easily deduced that the scale factor evolves as $a\\propto t^{1/2}$, the density as $\\rho\\propto t^{-2}$, and the temperature as  $T\\propto t^{-1/2}$. At $t=0$, the scale factor vanishes while the energy density and the temperature become infinite. This is why the point $t=0$ is known as the point of the initial cosmological singularity (Big Bang). As the universe expands, the density and the temperature decrease. When the universe is sufficiently cooled, particles and antiparticles annihilate each other, the photon-gas energy density falls off relatively rapidly and the main contribution of the matter density starts to come from the small excess of baryons over antibaryons, as well as from other fields and particles which now comprise the dark matter.  However, the hot universe theory suffers from a series of difficulties. The first difficulty is the singularity problem. The energy density goes to $+\\infty$ at $t=0$ and the solution cannot be formally continued for $t<0$.  A natural question is: ``What was before $t=0$?''. Of course, one can argue that space and time started with the Big Bang, so that this question is meaningless. However, this answer is not quite satisfactory. On the other hand, close to the Big Bang, the density can be arbitrarily large. This is in conflict with the laws of physics. When the density becomes high enough, quantum effects should be taken into account in the problem\\footnote{Remarkably, the question of the origin of space and time, and the importance of quantum mechanics in the early universe were discussed early by Lema\\^itre \\cite{lemaitrenature}: ``If the world has begun with a single quantum, the notion of space and time would altogether fail to have any meaning at the beginning''.}.  As a result, the Planck density $\\rho_P=c^5/G^2\\hbar=5.16\\, 10^{99}\\, {\\rm g/m^3}$, the Planck length $l_P=(G\\hbar/c^3)^{1/2}=1.62\\, 10^{-35}\\, {\\rm m}$, the Planck mass $M_P=(\\hbar c/G)^{1/2}=2.17\\, 10^{-5}\\, {\\rm g}$, the Planck temperature $T_P=M_P c^2/k_B=1.42\\, 10^{32}\\, {\\rm K}$, and the Planck time $t_P=(\\hbar G/c^5)^{1/2}=5.39\\, 10^{-44}\\, {\\rm s}$  should play a fundamental role in the very early universe.    The standard Big Bang theory suffers from other problems that are known as the flatness problem (or the fine-tuning problem), the horizon problem (or the causality problem), and the monopole problem. These problems can be solved at one stroke in the framework of inflationary cosmology \\cite{guth,linde}. In that scenario, the universe was in an unstable vacuum-like state described by some hypothetical scalar field $\\phi$ with its origin in quantum fluctuations. This  leads to an equation of state $p=-\\rho c^2$, implying a constant energy density $\\rho$, of the order of the Planck density $\\rho_P$, called the vacuum energy. As a result of this vacuum energy, the universe expands exponentially rapidly in a tiniest fraction of a second. This ``inflation'' can solve the above-mentioned problems. At the same time, the inflation theory predicts the scale invariant spectrum of the primordial density fluctuations and explains why the universe is spatially flat (or behaves {\\it as if} it were spatially flat due to the smallness of $k/a^2$). This is therefore a very attractive theory.    In this paper, we propose a model of the early universe based on a polytropic equation of state with index $n=1$. Before that, we first explain how we came to such a description.  In previous works \\cite{c1,c2,c3}, we explored the possibility that the universe is filled with self-gravitating Bose-Einstein condensates (BEC) with short-range interactions (see a detailed list of references in \\cite{c1} and recent additional references in \\cite{recent}). In the strong coupling limit (the so-called Thomas-Fermi approximation), the equation of state of the BEC is $p=(2\\pi a_s\\hbar^2/m^3)\\rho^2$, where $a_s$ is the scattering length \\cite{dalfovo}. This is the equation of state of a polytrope with index $n=1$ \\cite{chandra}. Initially, this model was introduced in order to describe dark matter halos. A nice feature of the BEC model is that it avoids the presence of cusps at the center of the halos  because of the effective repulsion due to the Heisenberg principle or because of the self-interaction of the bosons which plays the same role as the Pauli exclusion principle for fermions\\footnote{If dark matter halos are made of self-gravitating fermions, the cusps are prevented by the Pauli exclusion principle; see \\cite{c1} for discussions and references.}. Since cusps are not observed in real dark matter halos, the BEC model (or the fermion model) is favored as compared to the standard cold dark matter (CDM) model that predicts cuspy density profiles. Of course, at large scales, these models become equivalent.    The BEC model was then applied to cosmology  in order to describe the evolution of the universe as a whole \\cite{harko,harko2,c4,kl,vw,harko3}. In \\cite{c4}, we investigated the evolution of a ``BEC universe'' with an equation of state of the form $p=k\\rho^2c^2$. We considered the case of repulsive ($k>0$) and attractive ($k<0$) self-interaction.  More generally, we studied a ``cosmic fluid'' described by an equation of state of the form $p=(\\alpha \\rho +k \\rho^2)c^2$. When $k=0$, we recover the standard linear equation of state $p=\\alpha \\rho c^2$ describing radiation ($\\alpha=1/3$), pressureless matter ($\\alpha=0$), stiff matter ($\\alpha=1$), and vacuum energy ($\\alpha=-1$). Therefore, the generalized equation of state $p=(\\alpha \\rho +k\\rho^2)c^2$ can be viewed as the sum of a linear term describing a ``classical'' universe and a quadratic term due to the BEC. At late times, when the density is low, the effect of the BEC is negligible and we recover the classical universe. On the contrary, the effect of the BEC becomes important in the early universe where the density is high. Therefore, it is interesting to study how the BEC  affects the early evolution of the universe. In our previous paper \\cite{c4}, we assumed that BECs form after the radiation era, when the temperature has sufficiently decreased. This prevented us from extrapolating the solution before the radiation era (see the Remark in Sec. 8.2 of \\cite{c4}). In the present paper, we explore the possibility that the equation of state   $p=k\\rho^2c^2$ may hold {\\it before} the radiation era. Therefore, we assume that the early universe is described by an equation of state of the form $p=(\\alpha \\rho +k \\rho^2)c^2$ with $\\alpha=1/3$. Actually, the study of this generalized equation of state is interesting even if its origin is not connected to BECs. Furthermore, for this class of models, the Friedmann equations can be solved analytically, which is an additional source of motivation. First results were given in \\cite{harko,c4} where it was shown that the repulsive and attractive models behave very differently. For repulsive self-interactions ($k>0$), the universe starts at $t=0$ from a singularity at which the density is infinite but the radius finite. For attractive self-interactions ($k<0$), the universe  always existed in the past. For $t\\rightarrow -\\infty$, its density tends to a constant value while its radius tends to zero exponentially rapidly.  In the present paper, we complete the study of the generalized equation of state $p=(\\alpha \\rho +k \\rho^2)c^2$. First, we analyze its thermodynamical properties and study the temporal evolution of the temperature. Secondly, for $k<0$, we show that the exponential growth of the scale factor can  account for a phase of inflation with constant density that we identify with the Planck density $\\rho_P$. Finally, taking $\\alpha=1/3$, we study the transition between a pre-radiation era and the radiation era. We mention the analogy with a second order phase transition where the control parameter is the time $t$, the order parameter is the scale factor $a(t)$, and the Planck constant $\\hbar$ plays the role of finite size effects (the standard Big Bang theory is recovered for $\\hbar=0$). For $k<0$, the radius of the universe passes from $a_i=l_P=1.62\\, 10^{-35}\\, {\\rm m}$ at $t=t_i=0$ to $a_1=2.61\\, 10^{-6}\\, {\\rm m}$ at $t=t_1=1.25\\, 10^{-42}\\, {\\rm s}$ before entering in the radiation era and increasing algebraically as $t^{1/2}$. During the inflationary phase, the density remains approximately constant with the Planck value ($\\rho\\simeq \\rho_P=5.16\\, 10^{99}\\, {\\rm g}/{\\rm m}^3$) before decreasing as $t^{-2}$. The temperature passes from $T_i=5.54\\, 10^{-173}\\, {\\rm K}$ at $t=t_i=0$ to $T_1=3.93\\, 10^{31}\\, {\\rm K}$ at $t=t_1$ before decreasing algebraically as $\\sim t^{-1/2}$. It achieves its maximum value $T_e=7.40\\, 10^{31}\\, {\\rm K}$ at $t=t_e=1.27\\, 10^{-42}\\, {\\rm s}$.  For $k<0$, the initial radius of the universe $a(0)=a_1=2.61\\, 10^{-6}\\, {\\rm m}$ is already ``large'' so there is no inflation. The velocity of sound is less than the speed of light only for $t>t_i=8.32\\, 10^{-45}\\, {\\rm s}$, corresponding to $a_i=3.90\\, 10^{-6}\\, {\\rm m}$, $\\rho_i=1.29\\, 10^{99}\\, {\\rm g}/{\\rm m}^3$, and $T_i=1.64\\, 10^{32}\\, {\\rm K}$. This marks the beginning of the physical universe in this model.  The paper is organized as follows. In Sec. \\ref{sec_basic}, we recall the basic equations of cosmology and discuss famous cosmological models.  In Secs. \\ref{sec_ges} and \\ref{sec_dark}, we study the generalized equation of state  $p=(\\alpha \\rho +k\\rho^{1+1/n})c^2$  for any value of the parameters $\\alpha$, $k$ and $n>0$. We discuss two classes of solutions depending whether $k$ is positive or negative. In Secs. \\ref{sec_inf} and \\ref{sec_ni}, we consider a specific model of physical interest corresponding to $\\alpha=1/3$, $n=1$, and $k=\\pm 4/(3\\rho_P)$.  For $k<0$ (negative polytropic pressure), we get a model of inflationary universe without singularity. It describes in a unified manner the transition from a pre-radiation era to the radiation era. For $k>0$ (positive polytropic pressure), we get a model of non-inflationary universe with a new type of initial singularity. In Sec. \\ref{sec_analogy}, we develop an analogy with second order phase transitions where the Planck constant plays the role of finite size effects. ",
        "conclusions": "In this paper, we have carried out an exhaustive study of the generalized equation of state (\\ref{ges0}) for positive indices $n>0$. For $\\alpha=1/3$, $n=1$ and $k=-4/(3\\rho_P)$, this equation of state describes in a unified manner the transition between the pre-radiation era ($\\rho=\\rho_P$) and the radiation era ($\\rho\\propto a^{-4}$) in the early universe. It provides a model of early inflation. The case of negative indices $n<0$ is treated in Paper II. For $\\alpha=0$, $n=-1$ and $k=-\\rho_{\\Lambda}$, where $\\rho_{\\Lambda}$ is the cosmological density, this equation of state describes in a unified manner the transition between the matter era ($\\rho\\propto a^{-3}$) and the dark energy era  ($\\rho=\\rho_{\\Lambda}$) in the late universe. It provides a model of late inflation. Combining these two approaches, we obtain a simple non-singular model for the whole evolution of the universe based on two symmetric polytropic equations of state (see Paper II).   When we started our study, we were mainly interested in the equation of state $p=(\\alpha \\rho+k\\rho^2)c^2$ introduced in our previous paper \\cite{c4}. This equation of state was obtained in the context of  Bose-Einstein condensates, assuming that they are the constituents of  dark matter. Although BECs should form after the radiation era when the universe has cooled sufficiently, we wondered whether this equation of state could have some sense before the radiation era. In that picture, the pre-radiation era would correspond to primordial BECs with attractive self-interaction $\\lambda\\sim -(m/M_P)^4$ (see Appendix \\ref{sec_annbec}). Since the pre-radiation era is extremely cold (see Sec. \\ref{sec_inf}), the BEC model (which assumes $T=0$) may be an interesting suggestion to develop further.  However, we progressively realized that the study of a generalized equation of state of the form $p=(\\alpha \\rho+k\\rho^{1+1/n})c^2$ was of interest beyond the context of BECs. Therefore, our present point of view is more general. We believe that the equation of state (\\ref{ges0}) may have several origins that should be discussed in future works. By studying this equation of state in full generality, we realized that the positive indices $n>0$ describe the early universe while the negative indices $n<0$ describe the late universe. Furthermore, a positive polytropic pressure ($k>0$) leads to past or future singularities (or peculiarities) while a negative polytropic pressure ($k<0$) leads to non-singular models. They exhibit phases of early and late inflation associated with a maximum density $\\rho_{max}=\\rho_P$ (Planck density) corresponding to the vacuum energy in the past and a minimum density $\\rho_{min}=\\rho_{\\Lambda}$ (cosmological density) corresponding to the dark energy in the future.  Therefore, in the polytropic model, the description of the early and late universe appears to be very ``symmetric''. This result is obtained in a purely theoretical manner, without reference to observations. Strikingly, this symmetry is consistent with what we know about the real universe. This point will be further developed in Paper II.   Another motivation of our study was to discuss some aspects of the history of science and publicize (in the same vein as \\cite{luminet,nb}) the important contributions of Lema\\^itre in cosmology, which may not be sufficiently recognized. They include: 1. The first correct understanding of the mysterious de Sitter model \\cite{lemaitre1925}; 2. The re-discovery of the Friedmann equations and their thermodynamical interpretation \\cite{lemaitre1927}; 3. The instability of the Einstein static universe \\cite{lemaitre1927}; 4. The interpretation of the redshifts as a results of the expansion of the universe in accord with the Einstein theory of relativity \\cite{lemaitre1927};  5. The discovery of the Hubble law and the calculation of the Hubble constant two years before Hubble \\cite{lemaitre1927}; 6. The basis of Newtonian cosmology \\cite{lemaitreNewton,lemaitre1934}; 7. The correct place of the cosmological constant in the Poisson equation \\cite{lemaitrecosmo}; 8. The importance of the cosmological constant and its interpretation in terms of a vacuum energy density \\cite{lemaitre1934}; 9. The notion of a primordial singularity (primeval atom) that finally led to the Big Bang theory \\cite{lemaitresingular,lemaitreNewton,lemaitre1933}; 10. The importance of quantum mechanics in the early universe \\cite{lemaitrenature}.  Some other important contributions of Lema\\^itre have been collected recently in \\cite{luminetnew}.       \\appendix"
    },
    "1208/1208.2271_arXiv.txt": {
        "abstract": "Far-ultraviolet (FUV) radiation plays an important role in determining chemical abundances in protoplanetary disks.  \\ion{H}{1} Lyman $\\alpha$  is suspected to be the dominant component of the FUV emission from Classical T Tauri Stars (CTTSs), but is difficult to measure directly due to circumstellar and interstellar \\ion{H}{1} absorption. To better characterize the intrinsic Lyman $\\alpha$ radiation, we present FUV spectra of 14 CTTSs taken with the \\emph{Hubble Space Telescope} COS and STIS instruments.  H$_2$ fluorescence, commonly seen in the spectra of CTTSs, is excited by Lyman $\\alpha$ photons, providing an indirect measure of the Lyman $\\alpha$ flux incident upon the warm disk surface.  We use observed H$_2$ progression fluxes to reconstruct the CTTS Lyman $\\alpha$ profiles.  The Lyman $\\alpha$ flux correlates with total measured FUV flux, in agreement with an accretion-related source of FUV emission. With a geometry-independent analysis, we confirm that in accreting T Tauri systems Lyman $\\alpha$ radiation dominates the  FUV flux ($\\sim$1150 \\AA\\ - 1700 \\AA).  In the systems surveyed this one line comprises 70 - 90 \\% of the total FUV  flux. ",
        "introduction": "One of the most influential factors in determining the local composition and physical state of the gas in protoplanetary disks is the ultraviolet (UV) radiation field. Strong UV radiation is produced from hot gas in the magnetically active atmospheres of the central star, and at the magnetospheric shock where disk material collides with the stellar atmosphere \\citep{GSchmitt2008,Sacco2008}. UV photons have a profound influence on the gas heating \\citep{Jonkheid2004,Nomura2007,Woitke2009} and disk gas chemistry \\citep{Aikawa1999,Bethell2009,Fogel2011}. Surveys of the UV radiation field have been able to characterize the general spectral structure of the radiation \\citep{Valenti2000,Yang2012}. \\ion{H}{1} Lyman $\\alpha$ (Ly$\\alpha$) is a significant contributor to the UV radiation, comprising as much as 80\\% of the far-UV (FUV) emission produced in the stellar atmosphere and accretion shock \\citep{Bergin2003,Herczeg2004}. While the Ly$\\alpha$ radiation is intense close to the star, atomic hydrogen in the disk atmosphere at larger radii isotropically scatters Ly$\\alpha$ photons which provides greater penetration into the molecular disk.  This is quite different from UV photons at other wavelengths that are scattered solely by dust grains, leading to a Ly$\\alpha$ dominated radiation field - if Ly$\\alpha$ radiation is present  \\citep{Herczeg2005b,Bethell2011b}.  This is important because Ly$\\alpha$ dissociates many molecules such as H$_2$O and CH$_4$, and can lead to an overabundance of certain species, especially CN \\citep{Bergin2003,VanZadelhoff2003,VanDishoeck2006}.  While many strong UV transitions such as \\ion{C}{4} and \\ion{He}{2}  have been studied \\citep{Ingleby2011b,Yang2012}, an intrinsic Ly$\\alpha$ profile has only been well characterized in a single Classical T Tauri Star (CTTS) - TW Hya \\citep{Herczeg2004}. Due to circumstellar and interstellar \\ion{H}{1} absorption, it is impossible to directly measure the intrinsic Ly$\\alpha$ emission of most CTTSs. On the other hand, H$_2$ emission resulting from  Ly$\\alpha$ photoexcitation is prevalent throughout the FUV bandpass, providing an indirect probe of the Ly$\\alpha$.  Several Lyman and Werner band transitions reside at wavelengths coincident with the Ly$\\alpha$ profile, and have a relatively simple radiative transfer. These lines have been identified in previous observations of CTTSs \\citep{Brown1981,Valenti2000,Ardila2002,Herczeg2002,Herczeg2006,France2011b}. H$_2$ emission therefore provides an opportunity to study the circumstellar region where the Ly$\\alpha$ radiation interacts with the disk molecular gas, as well as characterize the strength of this key component of the FUV radiation field in T Tauri systems.  The impact of strong Ly$\\alpha$ radiation on disks has only recently been explored in detailed UV radiation transfer models \\citep{Bethell2011b} and chemical models \\citep{Fogel2011}.  Given the influence of these photons on gas throughout the disk, quantifying the Ly$\\alpha$ emission in CTTSs is crucial to produce accurate disk thermal/chemistry models.  In this letter, we reconstruct the Ly$\\alpha$ profile from observed H$_2$ emission lines in the FUV spectra of CTTSs. ",
        "conclusions": "We have demonstrated that Ly$\\alpha$ dominates the FUV emission of CTTSs. Ly$\\alpha$-fluoresced H$_2$ emission lines appear throughout the spectra of each CTTS in our survey. Circumstellar and interstellar \\ion{H}{1} partially attenuates the line centers of the Ly$\\alpha$ profiles in most targets and completely absorbs them in several targets.  The $observed$ Ly$\\alpha$ fluxes do not correlate with either the total H$_2$ emission or the summed FUV continuum flux. The strongest fluorescent H$_2$ progressions are used to reconstruct the Ly$\\alpha$ profile incident on the molecular disk, yielding $F_{Ly\\alpha}$ values accurate to within $\\sim20$\\% in most targets.   This uncertainty is caused by the distribution of possible \\ion{H}{1} and H$_2$ properties.  Our model assumptions (such as a single Gaussian emission component or  simple outflow absorber) may hide larger systematic errors, however future work will study these effects in a more detailed model.  We find that the intrinsic Ly$\\alpha$ comprises 81$\\pm9$\\% of the total FUV emission from CTTSs, compared with a fraction of only 15$^{+21}_{-15}$\\% for the observed Ly$\\alpha$ profiles. This demonstrates the need for Ly$\\alpha$ reconstruction to achieve accurate disk models of CTTSs.  It is clear from our results that the detection of the strong Ly$\\alpha$ line in TW Hya was not limited to a single object.  Rather, strong Ly$\\alpha$ emission dominates the FUV flux from all accreting (classical) T Tauri stars. Our measurements of the relative Ly$\\alpha$/FUV continuum flux only compare Ly$\\alpha$ to the FUV flux from 1160 - 1695 \\AA.  Thus if there is significant UV flux shortward of 1160 \\AA\\ then the strength of Ly$\\alpha$ relative to the FUV radiation could go down.  However, in TW Hya the flux below this limit is only 5\\% of the FUV flux below 1700 \\AA.  Similarly, \\citet{France2011a} have shown that the FUV continuum decreases to shorter wavelengths across the FUV bandpass. In addition, the absorption properties of grains strongly limit the propagation of UV photons near the Lyman limit.  The derived Ly$\\alpha$ fractions confirm the dominance of Ly$\\alpha$ in the FUV spectrum of the accreting young stars with disks.  This is important because Ly$\\alpha$ photons from the star will see the atomic hydrogen layer, with shallow angle of incidence, above the molecular surface.  Isotropic scattering will lead a significant fraction of the Ly$\\alpha$ flux to rain down on the disk with greater penetrating power than typical UV continuum photons \\citep{Bethell2011b}.  In addition, the Ly$\\alpha$ emission reprocessed by H$_2$ (which is scattered throughout the FUV spectrum) will also emit directly on the disk surface \\citep{France2012}.  These two effects increase the penetration of UV photons beyond the simple case where one assumes the UV photons observe the disk surface with shallow angle of incidence and with the propagation solely influenced by grains. In general this should lead to greater heating and additional chemical effects deeper in the disk system.  \\begin{figure}% \\centering \\includegraphics[width=9cm]{f3.pdf} \\caption{Observed (top) and model (bottom) Ly$\\alpha$ fraction (assuming $\\eta$=1) vs. total H$_2$ flux.  The diamond in each plot designates the average $f_{Ly\\alpha}$ value for the observed and model profiles.  The significant increase in Ly$\\alpha$ fraction from the observed to the model profiles demonstrates the importance of Ly$\\alpha$ reconstruction.} \\label{fig:Lyman_cont_frac} \\end{figure}   ES and KF thank Brian Wood for input on Ly$\\alpha$ profile reconstruction. This work was supported by NASA grants NNX08AC146 and NAS5-98043 to the University of Colorado at Boulder ($HST$ programs 11533 and 12036) and made use of data from $HST$ GO programs 8041 and 11616."
    },
    "1208/1208.0042_arXiv.txt": {
        "abstract": "If the dark matter of our galaxy is composed of heavy nuggets of quarks or antiquarks in a colour superconducting phase there will be a small but non-zero flux of these objects through the Earth's atmosphere. A nugget of quark matter will deposit only a small fraction of its kinetic energy in the atmosphere and is unlikely to be detectable. If however the impacting object is composed of antiquarks the energy deposited can be quite large and contain a significant charged particle content. These relativistic secondary particles will subsequently be deflected by the earth's magnetic field resulting in the emission of synchrotron radiation. This work argues that this radiation, along with a thermal component emitted from the nugget's surface, should be detectable at radio frequencies and that both present and proposed experiments are likely to prove capable of  detecting such a signal. ",
        "introduction": "\\label{sec:intro} Several recent experiments have  attempted to detect the radio emission associated with extensive air showers initiated by the impact of an ultrahigh energy cosmic ray on the earth's atmosphere \\cite{Ardouin:2009zp},\\cite{Huege:2012vk}, \\cite{Kelley:2012as}. This radio wavelength emission is generated by the deflection of secondary particles by the earth's magnetic field which produces synchrotron radiation. It is the purpose of this paper to demonstrate that experiments of this type will also capable of placing limits on dark matter in the form of heavy  quark matter nuggets. In the following section I offer a brief overview of the quark nugget dark matter model and it's observational consequences for both galactic observations and ground based detectors. Section \\ref{sec:radio_emis} extends the previous analysis of the air shower induced by a quark nugget passing through the earth's atmosphere \\cite{Lawson:2010uz} and estimates the synchrotron signal generated by such an event. Sections \\ref{sec:Efield_mag} and \\ref{sec:intens} then determine specific observable properties of the radio emission. Finally, section \\ref{sec:detect} offers a brief description of fututre detection prospects. \\subsection{dark matter as compact composite objects} \\label{subsec:gal} The microscopic nature of the dark matter is not yet established. While the majority of dark matter models introduce a new particle which is fundamentally weak in its interactions with visible matter it is only the  the dark matter  interaction cross section to mass ratio ($\\sigma / M$) that is observationally constrained. As such, a sufficiently heavy dark matter candidate may avoid observational constraints despite having a relatively large interaction strength. This is possible due to the fact that only the dark matter mass density is directly measured, a heavier dark matter candidate will have a lower number density and thus smaller flux through any detector. In the model to be discussed here the dark matter takes the form of heavy compact composite objects with nuclear scale densities and composed of the standard light quarks \\cite{Zhitnitsky:2002qa},\\cite{Oaknin:2003uv}. Several versions of this model have been proposed dating back to objects such as stranglets \\cite{Witten:1984rs}. Current observational constraints on the dark matter $\\sigma / M$ ratio require that these quark nuggets carry a mean baryonic charge of at least $10^{23}$. This lower bound is taken from a number of direct detection experiments as detailed in \\cite{Zhitnitsky:2006tu}. There is also a region of parameter space at $B\\sim 10^{29}  - 10^{31}$ likely excluded based on lunar seismic data \\cite{Herrin:2005kb}.  Quark nuggets may  be formed at the \\textsc{qcd} phase transition and remain stable over the lifetime of the universe. Nugget formation allows for the creation of nuggets of both matter and antimatter and may explain both the dark matter and the matter-antimatter asymmetry of the universe. This is possible if the production of antinuggets is favoured by a factor of $\\sim 1.5$ over the production of matter nuggets leading to dark matter which consists of two parts matter nuggets to three parts antimatter nuggets. The excess matter (that  not bound in the nuggets) then forms the visible matter of the universe in the observed five to one matter to dark matter ratio.  Originally proposed to explain the baryon asymmetry rather than any specific galactic observation this  model has been found to have several observational consequences for the galactic spectrum.  ${\\bullet}$ The quark nuggets are surrounded by an ``electrosphere\" of positrons the outer layers of which are bound with energies at typical atomic scales. The annihilation of these positrons with the electrons of the interstellar medium will result in a positronium decay line (and associated three photon continuum) in regions where both the visible and dark matter densities are large. In particular one should expect a $511keV$ line from the galactic centre \\cite{Oaknin:2004mn} \\cite{Zhitnitsky:2006tu}. Such a spectral feature is in fact observed and has been studied by the \\textsc{Integral} observatory \\cite{Jean:2005af}.  ${\\bullet}$ Positrons closer to the quark matter surface can carry energies up to the nuclear scale. If a galactic electron is able to penetrate to a sufficiently large depth it will no longer produce  the characteristic positronium decay spectrum but a direct $e^-e^+ \\rightarrow 2\\gamma$ emission spectrum \\cite{Lawson:2007kp}. Precisely modeling the transition between these two regimes allows for the determination of the strength of the MeV scale emissions relative to that of the 511keV line \\cite{Forbes:2009wg}. Observations by the \\textsc{Comptel} satellite show an excess above the gamma ray background predicted from galactic sources at the energy and intensity predicted by this model \\cite{Strong:2004de}.  ${\\bullet}$ Galactic protons may also annihilate with the antimatter comprising the quark nuggets. The annihilation of a proton within a quark nugget will produce hadronic jets which cascade down into lighter modes of the quark matter. If the energy in one of these jets reaches the quark surface it will excite the most weakly bound positron states near the surface. These excited positrons rapidly lose energy to the strong electric field near the quark surface. This process results in the emission of Bremsstrahlung photons at x-ray energies \\cite{Forbes:2006ba}. Observations by the \\textsc{Chandra} observatory indicate an excess in x-ray emissions from the galactic centre in precisely the energy range predicted \\cite{Muno:2004bs}.  ${\\bullet}$ The annihilation of visible matter within the nuggets heats them above the background temperature. The thermal spectrum from the nuggets may be predicted based on the emission properties of the electrosphere along with the annihilation rate at various positions within the galaxy \\cite{Forbes:2008uf}. The majority of this thermal energy is emitted at the eV scale where it is very difficult to observe against the galactic background. However the emission spectrum will extend down to the microwave scale where it may be responsible for the ``\\textsc{wmap} haze\" \\cite{Finkbeiner:2003im}.  ${\\bullet}$ The same thermal emission that may contribute to the \\textsc{wmap} haze associated with the galactic plane must also have been produced in the high density early universe when the isotropic matter density was comparable to that presently seen in  high density regions such as the galactic centre. Taking the nugget temperature scale from the galactic centre and scaling it with the cosmological expansion one predicts an excess in the isotropic radio background beginning a few decades in frequency below the \\textsc{cmb} peak \\cite{Lawson:2012zu}. The \\textsc{Arcade2} experiment show just such a low frequency isotropic radio excess \\cite{Fixsen:2009xn} which does not currently have an obvious astrophysical source.  The emission mechanisms involved in describing the full nugget spectrum span a very large energy range  from the modified thermal spectrum in the microwave up to energies associated with nuclear annihilations observable as gamma rays.  However, they make no significant contribution to the galactic spectrum at the GeV scale or above, an energy range across which the FERMI telescope has placed significant constraints on a possible dark matter contribution \\cite{Ackermann:2011wa}, \\cite{Ackermann:2010rg}, \\cite{Abdo:2010dk}.  While the uncertainties in the astrophysical backgrounds involved make an exact determination impossible a best fit to the galactic spectrum favours quark nuggets with a mean baryon number of roughly $B\\sim 10^{25}$. \\subsection{quark matter in the atmosphere} While the low number density of quark nuggets required to explain the observed dark matter mass density implies a flux many orders of magnitude below conventional dark matter models they  may still have observational consequences. Assuming the nuggets to have roughly nuclear scale densities the observational constraints on baryon number cited above imply a minimum mass of a gram and a radius of $\\sim 10^{-7}m$. Assuming the local dark matter mass density near the galactic plane average ($\\rho \\sim 1$Gev/cm$^3$) and a mean velocity at the galactic scale $v_g \\sim 200$km/s one may make an order of magnitude estimate the yearly flux of nuggets through the atmosphere as, \\begin{equation} \\label{eq:flux} \\frac{dn}{dt ~ dA} = n_N v_g = \\frac{\\rho_{DM}}{M_N} v_g = \\left( \\frac{B}{10^{25}} \\right) km^{-2} yr^{-1} \\end{equation} this flux is at a level comparable to that of ultra high energy cosmic rays near the \\textsc{GZK} cuttoff \\cite{Abraham:2008ru}. Consequently the current generation of large scale cosmic ray observatories have a sufficient collection area to conduct a meaning search for these objects provided their mass is near the lower end of the allowed range. Unlike standard cosmic rays the flux of quark nuggets will experience an annual variation as the earth moves through the background of galactic matter. This annual variation has been used in dark matter detection experiments such as \\textsc{Dama} \\cite{Belli:1999nz}.  In the case of nuggets of quark matter all energy deposited in the atmosphere comes through collisional momentum transfer. Relatively little energy is deposited and the nugget continues through the atmosphere with virtually no change in velocity. The observational prospect for such an event are very low. If however the nugget is composed of antiquarks (as the majority will be in the model under consideration) atmospheric molecules will annihilate on contact with the quark matter surface resulting in the release of substantial amounts of energy into the surrounding atmosphere. While the total annihilation of a quark nugget with $B=10^{25}$ would release $10^{14}$J of energy only a very small fraction of the nugget will annihilate in the time it takes to cross the earth's surface. As the nuggets carry sufficient momentum to traverse the entire atmosphere the limiting factor is the amount of atmospheric mass which they encounter.  This energy will be deposited in the form of thermal radiation from the nugget as well as relativistic particles and high energy gamma rays emitted in the nuclear interactions. The prospects for direct detection of secondary particles was discussed in \\cite{Lawson:2010uz}. This work focusses instead on the prospect of radio frequency detection. Both the thermal spectrum and the geo-synchrotron emission from the emitted particles will contribute to the radio spectrum, the magnitude and relative strength of these two emission mechanisms  will be approximated in section \\ref{sec:radio_emis}.  The total scale of the radio band spectrum is determined by the rate at which atmospheric molecules are annihilated within the nugget. This rate was estimated in \\cite{Lawson:2010uz} where it was shown to increase exponentially with the growing atmospheric density until the point at which the thermal energy produced by annihilations is sufficient to deflect further incoming matter. At this point the annihilation rate reaches an equilibrium point and does not increase further. The exact value of the equilibrium rate depends on details of emission from the quark surface but should occur when the surface temperature is on the order of $10keV$. For this temperature scale the annihilation rate saturates a few kilometers above the earth's surface. Above this height the annihilation rate is simply determined by atmospheric density and the physical cross section of the nugget. \\begin{eqnarray} \\label{eq:an_rate} \\Gamma_{an}&=& \\sigma_{N} v_{N} n_{at}(h) ~~~ h > h_{eq} \\\\ &=&  \\sigma_{N} v_{N} n_{at}(h_{eq}) ~~~ h < h_{eq} \\nonumber \\end{eqnarray} Where $\\sigma_N$ is the physical cross section of the nugget (on the order of $10^{-10} cm^2$ for the nugget mass range considered here), and $h_{eq}$ is the height at which the equilibrium annihilation rate is reached. ",
        "conclusions": "This paper has aimed to demonstrate the feasibility of detecting the presence of dark matter in the form of heavy nuggets of quark matter using radio frequency detectors. The passage of a quark nugget through the atmosphere will induce an extensive air showers involving many secondary charged particles which subsequently emit synchrotron radiation across the Mhz band as they are deflected by the earth's magnetic field. The resulting radio emission is likely to be detectable up to a few kilometers from the shower core. As such these events should be readily detectable by experiments intended to observe radio emission from ultra high energy cosmic rays.  While the intensity and event rate of these showers may be at a similar scale to that of air showers initiated by a single ultra high energy proton or nucleus antiquark nugget induced events have several easily observable distinguishing properties.  The nuggets require a time scale on the order of tens of milliseconds to traverse the earth's atmosphere and will generate a synchrotron signal over much of this time. As such the radio signal will be have a duration much longer than that of typical cosmic ray events which evolve on time scales orders of magnitude faster.  The nuggets carry galactic scale velocities and their flux will show a seasonal variation, however, any microscopic particle with so small a velocity will be insufficiently energetic to initiate an extensive air shower. Consequently, the detection of a seasonal variation in the air shower rate would be a strong indicator of a quark nugget contribution to this flux.  The composite nature of the primary particle in a quark nugget initiated air shower means that there is a thermal component to the spectrum in addition to the geosynchrotron emission generated by the secondary particles. While the thermal component is significantly lower than the synchrotron signal at the frequencies typically observed it's contribution increases at higher frequencies and may well be observable as distinct from the synchrotron signal.  If, as argued above, large scale cosmic ray detectors are also capable of observing the air showers induced by dark matter in the form of heavy nuggets of quark matter than these properties will allow the two components to be readily distinguished through observations at radio frequencies."
    },
    "1208/1208.5498_arXiv.txt": {
        "abstract": "Glycolaldehyde (HCOCH$_2$OH) is the simplest sugar and an important intermediate in the path toward forming more complex biologically relevant molecules. In this paper we present the first detection of 13~transitions of glycolaldehyde around a solar-type young star, through Atacama Large Millimeter Array (ALMA) observations of the Class~0 protostellar binary IRAS~16293-2422 at 220~GHz (6~transitions) and 690~GHz (7~transitions). The glycolaldehyde lines have their origin in warm (200--300~K) gas close to the individual components of the binary. Glycolaldehyde co-exists with its isomer, methyl formate (HCOOCH$_3$), which is a factor 10--15 more abundant toward the two sources. The data also show a tentative detection of ethylene glycol, the reduced alcohol of glycolaldehyde. In the 690~GHz data, the seven transitions predicted to have the highest optical depths based on modeling of the 220~GHz lines all show red-shifted absorption profiles toward one of the components in the binary (IRAS16293B) indicative of infall and emission at the systemic velocity offset from this by about 0.2$''$ (25~AU). We discuss the constraints on the chemical formation of glycolaldehyde and other organic species -- in particular, in the context of laboratory experiments of photochemistry of methanol-containing ices. The relative abundances appear to be consistent with UV photochemistry of a CH$_3$OH--CO mixed ice that has undergone mild heating.  The order of magnitude increase in line density in these early ALMA data illustrate its huge potential to reveal the full chemical complexity associated with the formation of solar system analogs. ",
        "introduction": "One of the most intriguing questions in studies of the chemistry of the early solar system is whether, how, when and where complex organic and potentially prebiotic molecules are formed. One of the key species in this context is glycolaldehyde (HCOCH$_2$OH). It is the simplest sugar and the first intermediate product in the formose reaction that begins with formaldehyde (H$_2$CO) and leads to the (catalyzed) formation of sugars and ultimately ribose, the backbone of RNA, under early Earth conditions \\citep[e.g.,][]{larralde95}. The presence of glycolaldehyde is therefore an important indication that the processes leading to biologically relevant molecules are taking place. However, the mechanism responsible for its formation in space is still unclear \\citep[see, e.g.,][]{woods12}.  Glycolaldehyde has so-far been detected in two places in space -- toward the Galactic center source SgrB2(N) (\\citealt{hollis00}; see also \\citealt{hollis01,hollis04,halfen06,requenatorres08}) and the high-mass hot molecular core G31.41+0.31 \\citep{beltran09}. Another compelling related discovery is that of ethylene glycol (``anti-freeze''; (CH$_2$OH)$_2$), the reduced alcohol variant of glycolaldehyde, found also toward SgrB2(N) at comparable abundances \\citep{hollis02}. Searches for glycolaldehyde in comets have so-far only resulted in upper limits, whereas ethylene glycol is detected toward Hale-Bopp and found to be at least 5 times more abundant than glycolaldehyde \\citep{crovisier04}. Comparisons between these species are therefore particularly interesting as their relative abundances potentially provide strong constraints on their formation and the chemical evolution from protostars to primitive solar system material.  The protostellar (Class~0) binary IRAS16293-2422 (IRAS16293 hereafter) at a distance of 120~pc \\citep{loinard08} has long been considered to be the best low-mass protostellar testbed for astrochemical studies \\citep[see, e.g.,][]{blake94,vandishoeck95,ceccarelli00a,schoeier02}, a status that has been further bolstered by the detection of a wealth complex organic molecules toward this source \\citep{cazaux03,caux11}. (Sub)millimeter wavelength interferometric studies of its chemistry have revealed strong differentiation among different species toward the two components in the binary (IRAS16293A and IRAS16293B; \\citealt{wootten89}) including complex organic molecules \\citep{bottinelli04iras16293,hotcorepaper,bisschop08,iras16293sma}.  With the Atacama Large Millimeter/submillimeter Array (ALMA) beginning operations a completely new opportunity has arisen for studies of the astrochemistry of solar-type stars. ALMA provides high sensitivity for faint lines, high spectral resolution which limits line confusion, and high angular resolution making it possible to study young stars on solar-system scales. In this letter we report the first potential discoveries of glycolaldehyde and ethylene glycol in a solar-type protostar from ALMA observations of IRAS~16293-2422. ",
        "conclusions": "\\label{chemistry} The similar spectral shapes of the lines of the complex organic molecules (in particular, the line widths and the absorption profiles toward IRAS16293B) as well as their similar spatial extent suggest that methyl formate, glycolaldehyde and other species coexist in the same gas. Furthermore, methyl formate and glycolaldehyde are fit by similar excitation temperatures in our simple LTE models.  Therefore, a discussion of the chemical relation between glycolaldehyde and methyl formate is relevant.  Our results indicate that methyl formate is a factor 10--15 more abundant than glycolaldehyde in the warm gas toward the two binary components of IRAS16293-2422. This value is consistent with previous measurements ranging from 52 in the hot core of SgrB2(N) \\citep{hollis01} and the upper limit of 34 in G34.41+0.31 \\citep{beltran09} to the average of 5--6.5 found on more extended scales toward SgrB2(N) \\citep{hollis00,requenatorres08}. Relative to the lower limit on the H$_2$ column density from the optically thick dust continuum emission toward IRAS16293B \\citep{chandler05} and taking into account the filling factor, the two species are estimated to have abundances relative to H$_2$ of 8$\\times 10^{-8}$ (methyl formate) and 6$\\times 10^{-9}$ (glycolaldehyde).  One of the major discussions concerning the origin of complex organic molecules in space is whether these form from second generation gas-phase reactions based on protonated CH$_3$OH (released from ices at high temperatures) or due to first generation reactions in the icy grain mantles, possibly induced by UV- or cosmic-ray irradiation \\citep[e.g.,][]{herbst09}. Grain-surface reactions are becoming increasingly more popular, especially since the formation of methyl formate through gas-phase reactions seems to be too inefficient to explain the observed abundances \\citep{horn04}.  \\cite{halfen06} compared a survey of the glycolaldehyde emission toward SgrB2(N) with formaldehyde (H$_2$CO), motivated by the first step in the formose reaction consisting of two H$_2$CO molecules combining to form HCOCH$_2$OH. Modeling the emission from H$_2$C$^{18}$O transitions detected in our SMA survey with the same excitation temperature as for methyl formate and glycolaldehyde suggests a formaldehyde to glycolaldehyde abundance ratio of 42--56 (using a $^{16}$O:$^{18}$O abundance ratio of 560 characteristic for the local ISM \\citep{wilson94}). This is in agreement with the estimate by \\cite{halfen06} of a ratio of 27 in SgrB2(N). \\citeauthor{halfen06} interpreted this ratio in favor of the formation of glycolaldehyde in space through a gas-phase formose reaction. However, as pointed out by \\cite{woods12} this and other gas-phase reactions tend to produce too little glycolaldehyde compared to observed abundances in an absolute sense.  Alternatively glycolaldehyde may be formed through grain-surface reactions in ices rich in methanol (CH$_3$OH) or its derivatives. From laboratory experiments \\cite{oberg09} found that photochemistry of UV-irradiated methanol ices mixed with significant amounts of CO can explain the observed fractions of oxygen-rich complex organics like methyl formate relative to methanol in sources such as IRAS16293 \\citep{bisschop08}. Although \\cite{oberg09} could not fully separate glycolaldehyde and methyl formate in their methanol-rich experiments, the derived lower limits on their abundances in CO-containing ices relative to methanol are 4\\% and 8\\%, respectively, consistent with the combined IRAS16293 results of \\cite{bisschop08} and those found here of $\\approx 1$\\% and 10--20\\%. Also, their upper limit of ethylene glycol relative to glycolaldehyde ($<$25\\%) agrees roughly with our ratio of 0.3--0.5. Experiments with irradiation of pure CH$_3$OH ices on the other hand produce too large ethylene glycol abundances relative to glycolaldehyde by an order of magnitude.  The CH$_3$OH:CO ratio is clearly critical: laboratory results for the CH$_3$OH:CO=1:10 mixtures show a 4 times larger abundance of glycolaldehyde relative to ethanol, whereas the ethanol lines detected here and in the previous SMA observations \\citep{bisschop08} suggests that ethanol is 3--5 times more abundant than glycolaldehyde in IRAS16293. As pointed out by \\cite{oberg09} the relative abundances of some species also strongly depend on the ice temperature, with the glycolaldehyde and ethylene glycol production requiring heating above $\\sim$30 K.  Thus, both exact ice composition (amount of CO mixed with CH$_3$OH) and temperature play a role in the chemistry; a combination of a moderately CO-rich ice and mild heating best reproduce the current data.  Still, additional systematic surveys of more species and sources are needed to constrain the surface formation mechanisms in more detail. These early data illustrate the enormous potential of ALMA for doing this.  The current sensitivity is already more than an order of magnitude better than that of previous single-dish or interferometric line surveys toward this source \\citep{caux11,iras16293sma}, revealing a line density toward IRAS16293B nearly ten times higher than before. Clearly, ALMA is posed to reveal many more complex organic molecules in young solar-system analogs."
    },
    "1208/1208.4564.txt": {
        "abstract": "We present multiwavelength data of the blazar 3C~454.3 obtained during an extremely bright outburst from November 2010 through January 2011. These include flux density measurements with the {\\it Herschel Space Observatory} at five submillimeter-wave and far-infrared bands, the {\\it Fermi} Large Area Telescope at $\\gamma$-ray energies, {\\it Swift} at X-ray, ultraviolet (UV), and optical frequencies, and the Submillimeter Array at 1.3 mm. From this dataset, we form a series of 52 spectral energy distributions (SEDs) spanning nearly two months that are unprecedented in time coverage and breadth of frequency. Discrete correlation anlaysis of the millimeter, far-infrared, and $\\gamma$-ray light curves show that the variations were essentially simultaneous, indicative of co-spatiality of the emission, at these wavebands. In contrast, differences in short-term fluctuations at various wavelengths imply the presence of inhomegeneities in physical conditions across the source. We locate the site of the outburst in the parsec-scale ``core'', whose flux density as measured on 7 mm Very Long Baseline Array images increased by 70\\% during the first five weeks of the outburst. Based on these considerations and guided by the SEDs, we propose a model in which turbulent plasma crosses a conical standing shock in the parsec-scale region of the jet. Here, the high-energy emission in the model is produced by inverse Compton scattering of seed photons supplied by either nonthermal radiation from a Mach disk, thermal emission from hot dust, or (for X-rays) synchrotron radiation from plasma that crosses the standing shock. For the two dates on which we fitted the model SED to the data, the model corresponds very well to the observations at all bands except at X-ray energies, where the spectrum is flatter than observed. ",
        "introduction": "The $\\gamma$-ray sky at high Galactic latitudes is dominated by highly variable active galactic nuclei termed ``blazars.'' Extreme $\\gamma$-ray apparent luminosities that can exceed $10^{48}$ erg s$^{-1}$ during outbursts in blazars are most easily understood as the result of Doppler boosting of the nonthermal emission from a relativistic plasma jet pointing within several degrees of our line of sight \\citep[e.g.][]{dermer95}. The Doppler effect also shortens the observed timescale of variability and accentuates the amplitude of the flux changes. The emission at radio to optical --- and in some cases up to X-ray --- frequencies is synchrotron radiation from ultra-relativistic electrons gyrating in magnetic fields inside the jet.  The electrons that give rise to the synchrotron emission also create X-rays and $\\gamma$-rays through inverse Compton scattering. The synchrotron radiation, as well as other types of emission from various regions in the galactic nucleus, can provide seed photons for scattering. Neither the main source(s) of these seed photons nor the processes that energize the radiating electrons have yet been identified with certainty. A promising method for doing so is to observe and model both the spectral energy distribution (SED) at various stages of an outburst and the time delays between variations at millimeter, infrared (IR), optical, ultraviolet (UV), X-ray, and $\\gamma$-ray bands during outbursts. The goal is to use the results of such an analysis to identify the location of the scattering regions (i.e., in the vicinity of the accretion disk $\\lesssim 10^{17}$ cm from the central supermassive black hole, within the inner parsec, or farther out) and, by doing so, infer the source of the seed photons [nonthermal emission from the jet or thermal emission from the accretion disk, broad emission-line region (BLR), or parsec-scale hot dust]. This information can then be used to define, or place strong constraints on, the physical conditions in the inner $\\sim 10$ pc of the jet (e.g., magnetic field, electron energy density, energy gains and losses of the relativistic electrons, and changes in bulk Lorentz factor).  The flat-radio-spectrum quasar 3C~454.3 (redshift of 0.859) is particularly well-suited for such a study. Analysis of time sequences of images with angular resolution of $\\sim 0.1$ milliarcseconds (mas) with the Very Long Baseline Array (VLBA) at a frequency of 43 GHz shows that this blazar contains a narrow (opening half-angle $\\lesssim 1^\\circ$) relativistic jet pointed within $2^\\circ$ of our line of sight \\citep{J05}. An extremely bright multiwavelength outburst in 2005 provided an example of the high-amplitude variability of the jet \\citep{Vil06,J10}.  On 2010 October 31, observers discovered that 3C~454.3 was undergoing a pronounced flare at near-IR wavelengths \\citep{carrasco2010}. The flare extended to millimeter, IR, optical, UV, X-ray, and $\\gamma$-ray bands  \\citep[see][and references therein]{vercellone2011}. In 2010 November, 3C~454.3 attained the highest flux ever observed at $\\gamma$-ray energies near 1 GeV \\citep{vercellone2010,vercellone2011,striani2010,sanchez2010,abdo2011}, peaking on November 19-20. Based on the extraordinary nature of the outburst, we obtained a series of {\\it Herschel} and {\\it Swift} target of opportunity observations of the quasar\\footnote{{\\it Herschel} is a European Space Agency space observatory with science instruments provided by European-led principal investigator consortia and with important participation from NASA.}. %changes made below This was the first time that {\\it far}-IR, X-ray, and $\\gamma$-ray space telescopes were all available to observe simultaneously during a truly extraordinary blazar outburst. We have previously observed flares of blazars at mid-IR and near-IR bands with the {\\it Spitzer Space Telescope}, first without and then with concurrent {\\it AGILE} $\\gamma$-ray observations \\citep[][ Wehrle et al., in preparation]{ogle2011}. In those observations, we unexpectedly discovered that there were two peaks in the synchrotron SED. The synchrotron SED of BL Lac has also been found to contain a double hump profile at higher frequencies \\citep{raiteri2010}. Such structure in the SED could imply the existence of two distinct sites of IR emission in the jet. One possibility is that both a relatively quiescent component --- e.g., the quasi-stationary ``core'' seen in millimeter-wave VLBA images --- and a propagating disturbance --- appearing as a superluminal knot in sequences of such images --- are present during outbursts. On the other hand, during a singularly bright flare a new knot should dominate the flux, producing a single peak in the SED.  Two peaks in the synchrotron SED can also be explained by two separate electron-positron populations. Hadronic proton-proton and proton-photon collisions could make enough pions to produce, via decays and electromagnetic cascading, a secondary electron-positron population. Alternatively, two electron acceleration mechanisms, e.g., magnetic reconnections and shocks, could operate, or  the physical conditions could vary across the source.  \\citet{J10} have linked strong $\\gamma$-ray flares with the appearance of superluminal knots at millimeter wavelengths in 3C~454.3. If this is the case for all such events, we expect to see increased activity in the jet in VLBA images at 7 mm that we obtained during and after the outburst.  Here we present extensive multiwavelength data during the outburst in 3C~454.3 and discuss its implications. Our data extend the results presented by \\citet{abdo2011} and \\citet{vercellone2011}, especially through the addition of {\\it Herschel} and VLBA observations. We also offer a different interpretation of the location of the event and the source of seed photons that were scattered up to $\\gamma$-ray energies. When translating angular to linear sizes, we adopt the current standard flat-spacetime cosmology, with Hubble constant $H_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\Lambda=0.73$. At a redshift $z=0.859$, 3C~454.3 has a luminosity distance $d_\\ell$ = 5.489 Gpc, and 1 mas corresponds to a projected distance of 7.7 pc in the rest frame of the host galaxy of the quasar. ",
        "conclusions": "The variations of flux at millimeter, far-IR, and $\\gamma$-ray bands during the outburst of 3C~454.3 from 2010 November to 2011 January were strongly correlated, with the $\\gamma$-ray variations lagging those at 1.3 mm and 160 $\\mu$m by $1.0\\pm 0.5$ days. Despite the strong correlations, the substructure of the outburst reveals significant differences at various wavebands. This suggests that physical substructure is present in the jet, which we interpret as evidence for turbulent processes.  The nearly simultaneous peak of the main flare at 1.3 mm and at $\\gamma$-ray energies implies that the flaring component of the multiwavelength emission was optically thin and dominated the emission at frequencies higher than 43 GHz. At 1.3 mm, the core region and knot K09 (ejected in late 2009) were the main emitters. Analysis of the light curve of the two components leads to the conclusion that the core region was responsible for the 1.3 mm outburst, while the flux of K09 declined during the event. The lack of a significant time delay between the flare at 1.3 mm and $\\gamma$-ray energies requires that the flare takes place in a region that is transparent at 1.3 mm. In support of this, the core region brightened by 70\\% at 7 mm during the early stages of the outburst bracketing the main flare. The data therefore favor a location of the outburst within the parsec-scale core rather than within the BLR. This conflicts with the model proposed by \\citet{vercellone2011}, who place the flare in a region $3.6\\times 10^{16}$ cm in radius only 0.05 pc from the central engine. In addition, later VLBA observations have identified a bright superluminal knot that was blended with the core during the outburst. The disturbance creating the flare therefore appears to have created the new knot.  Our {\\it Herschel} far-IR observations clearly define the low-frequency side of the maximum in the synchrotron SED, from which we infer that the SED peaked at a wavelength shorter than 70 $\\mu$m. Two weeks before the flare, IRTF observations located the peak at 11 $\\mu$m. The SED of the main flare did not shift toward lower frequencies as rapidly as expected in shock or expanding plasmoid models. This leads us to interpret the outburst in terms of a standing shock in the jet, across which turbulent plasma flows. A major increase in the energy density of the plasma causes an outburst as it crosses the shock, which has a conical structure. If the conical shock terminates in a transversely oriented Mach disk, a strong shock that decelerates and compresses the flow greatly, emission from the Mach disk can provide the main source of seed photons for inverse Compton scattering to generate the X-ray and $\\gamma$-ray emission. The TEMZ code  produces model SEDs, following this scenario, that match the millimeter to optical and $\\gamma$-ray spectra quite well although the observed X-ray spectrum is somewhat steeper than in the model calculations. Alternatively, if asymmetries in the jet flow prevent the formation of a Mach disk, thermal IR emission from hot dust can provide the requisite seed photons. However, the dust would need to have a luminosity $\\sim 1\\times 10^{46}$ erg s$^{-1}$ --- half that of the accretion disk --- and would need to be very patchy. % changed again below. In the TEMZ model, the volume filling factor of the emission is inversely related to the frequency of observation. Because of this, the average degree of linear polarization, as well as the level of variability of both the flux and polarization, should increase with frequency in a manner that is quantitatively related to the amount of steepening of the spectral slope toward high frequencies \\citep{mj10}. Another prediction of the model is that details of the flare profiles should change randomly from one flare to another.   Our observations of 3C~454.3 therefore present significant challenges to all theoretical models that have been proposed thus far. Since the polarization properties of blazars require irregularities in the magnetic field that are most easily understood as the effects of turbulence, scenarios similar to the TEMZ model appear to be necessary to reproduce the observed multi-waveband behavior of blazars. Whether this class of models can succeed in doing so requires further development of such models as well as more comprehensive observations of distinct outbursts in blazars."
    },
    "1208/1208.5517_arXiv.txt": {
        "abstract": "We report the discovery by the Wide-field Infrared Survey Explorer of the $z=2.452$ source \\longw1814, the first hyperluminous source found in the WISE survey. \\w1814\\ is also the prototype for an all-sky sample of $\\sim1000$ extremely luminous ``W1W2-dropouts\" (sources faint or undetected by WISE at 3.4 and $4.6\\ \\mu$m and well detected at 12 or $22\\ \\mu$m). The WISE data and a $350\\ \\mu$m detection give a minimum bolometric luminosity of $3.7\\times10^{13}L_\\odot$, with $\\sim10^{14}L_\\odot$ plausible. Followup images reveal four nearby sources: a QSO and two Lyman Break Galaxies (LBGs) at $z=2.45$, and an M dwarf star. The brighter LBG dominates the bolometric emission. Gravitational lensing is unlikely given the source locations and their different spectra and colors. The dominant LBG spectrum indicates a star formation rate $\\sim300 M_\\odot {\\rm yr}^{-1}$, accounting for $\\simlt10\\%$ of the bolometric luminosity. Strong $22\\ \\mu$m emission relative to $350\\ \\mu$m implies that warm dust contributes significantly to the luminosity, while cooler dust normally associated with starbursts is constrained by an upper limit at 1.1 mm. Radio emission is $\\sim10\\times$ above the far-infrared/radio correlation, indicating an active galactic nucleus is present. An obscured AGN combined with starburst and evolved stellar components can account for the observations. If the black hole mass follows the local $M_{\\rm BH}$-bulge mass relation, the implied Eddington ratio is $\\simgt4$. \\w1814\\ may be a heavily obscured object where the peak AGN activity occurred prior to the peak era of star formation. ",
        "introduction": "}   The Wide-field Infrared Survey Explorer (WISE) launched on 2009 Dec. 14, and began surveying the sky on 2010 Jan. 14, completing its first full coverage in July 2010.  WISE achieves $5 \\sigma$ point source sensitivities of better than 0.08, 0.11, 1, and 6 mJy at 3.4, 4.6, 12 and $22\\ \\mu$m (hereafter referred to as W1, W2, W3, and W4) in a single coverage on the ecliptic, consisting of 8 or more exposures at each sky location \\citep{Wright:10}. Sensitivity improves away from the ecliptic due to denser exposure overlap and lower zodiacal background. The survey continued in W1 and W2 after the cryogen was exhausted at the end of Sept. 2010, and concluded 2011 Feb. 1, having achieved two complete sky coverages in these two bands.  The WISE all-sky data release was issued on 2012 Mar. 14\\footnote{\\tt http://wise2.ipac.caltech.edu/docs/release/allsky}.  The primary science objectives for WISE are to identify the coldest and nearest brown dwarfs to the Sun \\citep[see e.g., ][]{Mainzer:11a, Cushing:11, Kirkpatrick:11}, and the most luminous, dusty, forming galaxies (Ultra-luminous Infrared Galaxies or ULIRGs). With regard to the most luminous objects, the dominant sources of energy production in the Universe are fusion in stars and gravitational accretion onto super-massive black holes. The tight correlation between the masses of super-massive black holes at the centers of galaxies, and the masses of the stellar bulges in these galaxies \\citep{Magorrian:98, Ferrarese:00,Gebhardt:00}, implies the two formation processes are intimately connected. Bulge stellar populations are old and quiescent today, and so must have formed in the distant past, when the cosmic star formation rate was much higher \\citep[e.g.,][]{Hopkins:04}. Similarly the peak era of accretion onto super-massive black holes was at redshifts $z\\sim2$, when luminous quasars were common \\citep[e.g.,][]{Richards:06a, Assef:11}. There is substantial evidence that dust absorbs much of the UV/optical luminosity generated by the formation of massive galaxies and their central black holes at $z > 1$ \\citep[e.g.,][]{Blain:04, LeFloch:05, Stern:05b, Hickox:07}. The dust is heated in the process, and most of the luminosity emerges at IR wavelengths, creating a ULIRG.  The most extreme examples of ULIRGs are therefore likely associated with the major formation events of the most massive galaxies. Such objects should appear in the sensitive infrared all-sky WISE survey, and may be missed in surveys with substantially smaller areas by e.g., \\spitzer\\ and {\\it Herschel}. Here we report on \\longw1814\\ (hereafter WISE 1814+3412), the first hyper-luminous infrared galaxy ($L_{\\rm IR} > 10^{13} L_\\odot$) discovered by WISE.  Magnitudes are converted to flux densities using zeropoint values of 3631 Jy for the AB system {\\it g'} and {\\it r'} bands.  Other magnitudes are on the Vega system, using zeropoints of 1594 and 666.7 Jy for the 2MASS system {\\it J} and ${\\it K_s}$ bands; 280.9 and 179.7 Jy for {\\it Spitzer} IRAC [3.6] and [4.5]; and 306.7, 170.7, 29.04 \\& 8.284 Jy for WISE W1 through W4 respectively \\citep{Wright:10}. Luminosities are calculated assuming $\\Omega_M = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70\\ {\\rm km\\ s}^{-1} {\\rm Mpc}^{-1}$. ",
        "conclusions": "With a minimum bolometric luminosity of $3.7 \\times 10^{13} L_\\odot$, and more likely $L_{\\rm bol} \\sim 9 \\times 10^{13} L_\\odot$, \\w1814\\ easily qualifies as a hyper-luminous infrared galaxy, the first identified by WISE.  Its W4 (22 $\\mu$m) flux density and redshift exceed that of the brightest $24\\ \\mu$m source DOG in the $8\\ {\\rm deg}^2$ Bo\\\"otes field \\citep[SST24 J1428+3541;][]{Desai:06}, giving it a rest $5\\ \\mu$m luminosity $3\\times$ higher. Its luminosity probably exceeds the estimated $4\\times 10^{13} L_\\odot$ of the most luminous obscured quasar in \\citet[SWIRE J164216.93+410127.8;][]{Polletta:08}, which has $4\\times$ lower flux at $24\\ \\mu$m. It approaches the $\\sim 2 \\times 10^{14} L_\\odot$ of the most luminous known quasars such as S5~0014+813 \\citep{Kuhr:83}, HS~1700+6416 \\citep{Reimers:89}, and SDSS~J074521.78+473436.2 \\citep{Schneider:05}.  Several W1W2-dropouts reported in \\citet{Wu:12} are even more luminous than \\w1814, with {\\em minimum} $L_{\\rm bol}$ up to $1.8\\times10^{14} L_\\odot$.  \\w1814\\ is also a massive, actively star forming LBG, with an extinction corrected UV star formation rate of $\\sim 300\\ M_\\odot {\\rm yr}^{-1}$ and a stellar mass of $\\sim 3 \\times 10^{11} M_\\odot$ \\citep[several times the stellar mass of a field $L*$ galaxy today;][]{Baldry:08}.  The specific star formation rate of $\\sim 1\\ {\\rm Gyr}^{-1}$ places \\w1814\\ in the starburst galaxy regime, but even the upper end of the extinction corrected UV SFR estimates account for $\\simlt 10\\%$ of the bolometric luminosity. With no AGN signatures in the UV spectrum, it is tempting to assume that additional dust-obscured star formation accounts for the high bolometric luminosity.  However, the upper limit at 1.1 mm would be greatly exceeded unless the star forming dust was significantly warmer than 35 K.  This possibility merits consideration. \\citet{Kovacs:06b} present evidence that starburst dominated sub-millimeter galaxies have higher temperatures at higher redshifts and luminosities. The SED of the extreme low metallicity dwarf starburst galaxy SBS 0335-052 peaks at $30 \\mu$m \\citep{Houck:04}, and \\citet{Hirashita:04} suggest that virtually all high redshift star formation may occur in the ``active mode\" for which SBS 0335-052 is the prototype.   The active mode is characterized by dense, compact star-forming regions in which dust from core collapse SNe warms as it screens the gas from UV photons, enabling runaway star formation as the gas continues to cool on a dynamical timescale. Yet it is not obvious that this mechanism can be scaled up to the levels required to produce $\\sim 10^{14} L_\\odot$, i.e. a SFR $\\sim 10^4 M_\\odot {\\rm yr}^{-1}$ and $\\sim 10^8$ O stars.  Since \\w1814\\ contains a modestly powerful radio source, an active galactic nucleus is almost certainly present. An obscured AGN powered by a super massive black hole (SMBH) could supply the luminosity while being consistent with the absence of AGN signatures in the UV spectrum. In fact, despite the clear starburst signature from the optical spectrum, the lack of submm emission and predominance of hot dust suggest that in \\w1814\\ \\citep[and other W1W2-dropouts;][]{Wu:12} we may be seeing an SED more like the ``intrinsic\" SED from the dusty central AGN region than is the case for many quasars \\citep[see Figure 6 of][]{Netzer:07}.  The formation and coevolution of AGN and stellar populations is the subject of much current research. Theoretical scenarios suggest the initial starburst rapidly enshrouds a forming galaxy in cold dust which emits at submm wavelengths. This is followed by an increase in fueling of the SMBH, triggering an AGN, generating warmer dust and mid-IR emission. Finally outflows from the AGN or starburst clear the dust and gas, revealing an optical quasar, and then removing the fuel which powers both starburst and quasar, leaving a quiescent massive galaxy \\citep[e.g.,][]{Sanders:88, Hopkins:06, Narayanan:09}.  The black hole mass needed to power the bolometric luminosity is $2.8 \\times 10^9 (L_{\\rm Edd}/L_{\\rm bol}) M_\\odot$. But using the local $L_{\\rm bulge}/M_{\\rm BH}$ relation, the SED fit stellar component absolute magnitude of $M_K = -26.28$ gives a black hole mass of only $1.4 \\times 10^9 M_\\odot$ \\citep{Graham:07}. If the relation is actually between SMBH mass and bulge {\\em mass} rather than bulge luminosity, as seems plausible, the discrepancy is larger. At $z=2.452$ the stellar mass for a given $K$ luminosity is likely less than half as large as today. Therefore, assuming the bolometric luminosity of the AGN template fit is correct, and that the \\citet{Graham:07} relation adjusted for lower $(M/L)_K$ at $z=2.45$ applies, the implied Eddington ratio is $\\simgt 4$. For comparison, \\citet{Kollmeier:06} measure a typical Eddington ratio in luminous quasars of 0.25, and \\citet{Kelly:10} find 0.1 to be representative.  If \\w1814\\ has a similar Eddington ratio, then one of these assumptions does not hold.  The local stellar mass to SMBH mass relation could apply if the bolometric luminosity is mainly from star formation with warm dust similar to that in low metallicity starburst dwarf galaxies, but scaled up a factor of a thousand. There is evidence that starbursts can dominate luminosities even in sources with spectral signatures of quasars \\citep[e.g.,][]{Netzer:07, Polletta:08}. Alternatively, the stellar mass may have been underestimated by a factor of several due to greater extinction at $4.5\\ \\mu$m,  but the extinction corrected stellar mass of $>10^{12} M_\\odot$ is then extreme. Note this argument holds for AGN with $L > 10^{14} L_\\odot$ in general, and the largest SMBH masses in the most luminous objects may be well above the normal relation \\citep[e.g.,][]{McConnell:12, Salviander:12}. Another possibility is that \\w1814\\ is in an unstable transitional stage where the Eddington ratio is well above 1.  Recently, \\citet{Kawakatu:11} have argued that SMBH's with super-Eddington accretion will be characterized by a low ratio of near-IR to bolometric luminosity, which is essentially the selection criterion for \\w1814.  We consider it more likely that \\w1814\\ lies well off the local $M_{BH}$-bulge mass relation, suggesting it may be in an early stage of its evolution where AGN activity is dominant and the object has yet to build the bulk of its stellar mass.  This would, however, contradict the expectation that AGN activity should lag the peak of star formation \\citep[see e.g.,][]{Hopkins:12}. In this context, the presence of 3 systems (components A, C, and D) at the same redshift and within 50 kpc suggests that a significant amount of stellar mass assembly lies in \\w1814's future."
    },
    "1208/1208.3107_arXiv.txt": {
        "abstract": "In order to determine the Galactic distribution of supernova remnants (SNRs) there are two main difficulties: (i) there are selection effects which mean that catalogues of SNRs are not complete, and (ii) distances are not available for most SNRs, so distance estimates from the `$\\Sigma{-}D$' relation are used. Here I compare the observed distribution of 69 `bright' SNRs with Galactic longitude with that expected from the projection of various model Galactocentric radius distributions. This does not require distances from the `$\\Sigma{-}D$' relation, and selecting only `bright' remnants aims to avoid major issues with the selection effects. Although this method does not provide a direct inversion to the 3-D distribution of SNRs in the Galaxy, it does provide useful constraints on the Galactocentric radius distribution. For a combined power-law/exponential model for SNR surface density variation with Galactocentric radius, the best fitted distributions are more concentrated towards lower radii than the distribution derived by \\citet{1998ApJ...504..761C}. ",
        "introduction": "Supernova remnants (SNRs) are important sources of energy and high energy particles in the Galaxy. Consequently the distribution of SNRs with Galactocentric radius is of interest for studies of cosmic rays in the Galaxy, and the high energy $\\gamma$-rays they produce from interaction with the interstellar medium (see, for example, \\citet{2012ApJ...750....3A} and \\citet{2012ApJ...752...68V}). Here I discuss some of the problems in constructing the Galactic distribution of SNRs, particularly due to selection effects, and the fact that distances are not available for most SNRs. I then compare the observed distribution of SNRs with Galactic longitude with the expected distribution from various (simple) models, in order to provide constraints on the Galactic distribution of SNRs. ",
        "conclusions": "Observational selection effects, and the lack of distances for many Galactic SNRs means that deriving the Galactic SNR distribution is not straightforward. For a sample of 69 `bright' SNRs -- i.e.\\ those not strongly affected by selection effects -- the observed $l$-distribution is shown not to be consistent with the Galactic distribution with Galactocentric radius derived by \\citet{1998ApJ...504..761C}.  \\begin{theacknowledgments}  I thank Irina Stefan for useful comments on early drafts of this paper.  \\end{theacknowledgments}"
    },
    "1208/1208.6514_arXiv.txt": {
        "abstract": "{The recent discovery by Gillessen and collaborators of a cloud of gas falling towards the Galactic Center on a highly eccentric orbit, diving nearly straight into the immediate neighborhood of the central supermassive black hole, raises the important question of its origin. Several models have already been proposed. Here we suggest that a recent nova outburst has ejected a ring-like shell of gas. Viewed at high inclination, that could account for the mass, head and tail structure, and the unusually high eccentricity of the observed cloud in a natural way, even as the nova moves on an orbit quite normal for the young stars in the close neighborhood of the Galactic Center. We illustrate this by calculating orbits for the head and tail parts of the ejecta and the nova that has produced it. We briefly discuss some of the questions that this model, if true, raises about the stellar environment close to the Galactic Center. } ",
        "introduction": "The discovery  by Gillessen et al. (2012) of a cloud of gas freely falling toward the Galactic center on an orbit diving deeply into the immediate neighborhood of the central supermassive black hole, is a great achievement that raises important questions for our understanding of the Galactic Center. The cloud was recognized as a moving object in L'-band observations. In the position-velocity maps of Br$\\gamma$ emission, obtained with SINFONI, a bright `head' of emission appears together with a `tail' of lower surface brightness (Gillessen et al. 2012, Fig.2). How did this cloud form and how did it get the peculiar orbit so close to the central supermassive black hole that it already now shows the destructive shear of gravity? The surprising answer might be: We are witnessing the recent outburst of a nova, the ejecta now freely falling toward the Galactic Center and being observed as they disperse in the close neighborhood of the black hole. ",
        "conclusions": "We have investigated the possibility that the cloud observed by Gillessen et al. (2012) in the close neighborhood of the Galactic Center results from a nova outburst. Mass, velocities, and the presence of dust agree with those of observed nova shells. We find that simple examples for the orbits of ejected matter allow interpretation of the cloud head in its observed peculiar orbit around the Galactic Center as part of the expanding shell of a nova that exploded around the year 2000. We note that different parts of a ring-like shell will appear at different brightness to the observer, the cloud head being the brightest one. The velocity with which the observed parts of the cloud structure, head and tail, move apart from each other, suggests that the tail is the, for the observer, second brightest part of the shell, ejected in the direction opposite to that of the head. The high eccentricity of the cloud head can be understood as caused by the addition of the velocities of the nova itself and the ejecta in that particular direction.  Though the nova origin model seems promising it still needs further evaluation. We note a wide range of possible parameters for ejected nova shells and their orientation with respect to the observer, which could allow a detailed comparison of such a model with the observations of that extraordinary cloud event."
    },
    "1208/1208.3870_arXiv.txt": {
        "abstract": "\\noindent In stars with $M_\\ast \\lesssim 2 \\msol$, nuclear burning of helium starts under degenerate conditions and, depending on the efficiency of neutrino cooling, more or less off-center. The behavior of the \\emph{centers} of low-mass stars undergoing core helium ignition on the $\\log\\rho - \\log T$ plane is not thoroughly explained in the textbooks on stellar evolution and the appropriate discussions remain scattered throughout the primary research literature. Therefore, in the following exposition we collect the available knowledge, we make use of computational data obtained with the open-source star-modeling package MESA, and we compare them with the results in the existing literature. The line of presentation follows essentially that of \\citet{Thomas1967} who was the first who outlined correctly the stellar behavior during the off-center helium flashes that lead to central helium burning. The exposition does not contain novel research results; it is intended to be a pedagogically oriented, edifying compilation of pertinent physical aspects which help to \\emph{understand} the nature of the stars. ",
        "introduction": "\\label{sect:modeling} All the data to which the ensuing discussion will refer were obtained with the versatile MESA code suite, which is comprehensively described in \\citet{Paxton2011}.  The computations were performed using release version 3251. Since we are only interested in the generic behavior of the model stars, we kept the prescription of the physical ingredients as simple as possible. Hence, except if mentioned explicitly otherwise in the text, we computed the stellar-evolution models without rotation, magnetic fields, and we neglected mass loss.  Convection was treated according to Henyey's MLT prescription and we adopted the Schwarzschild criterion for convective instability; the mixing-length was set ad hoc to 1.8 pressure scale-heights.  Even though different microphysics and in particular more appropriate treatment of dynamical convection will change the results quantitatively, probably mainly with respect to abundance profiles and the associated consequences, we are positive that the following discussion caught the generic nature of low-mass stars' evolution through the onset of core helium burning and~--~viewed on the HR diagram~--~their relatively fast transition from the top of the giant branch onto the horizontal branch or into the clump on the giant branch. ",
        "conclusions": ""
    },
    "1208/1208.3277_arXiv.txt": {
        "abstract": "It has long been speculated that the nature of the hard X-ray corona may be an important second driver of black hole state transitions, in addition to the mass accretion rate through the disk.  However, a clear physical picture of coronal changes has not yet emerged.  We present results from a systematic analysis of Rossi X-ray Timing Explorer observations of the stellar mass black hole binary \\j.  All spectra with significant hard X-ray detections were fit using a self-consistent, relativistically-blurred disk reflection model suited to high ionization regimes.  Importantly, we find evidence that both the spectral and timing properties of black hole states may be partially driven by the height of the X-ray corona above the disk, and related changes in how gravitational light bending affects the corona--disk interaction.  Specifically, the evolution of the power-law, thermal disk, and relativistically--convolved reflection components in our spectral analysis indicate that: (1) the disk inner radius remains constant at \\rin$=1.65\\pm0.08~GM/c^2$ (consistent with values found for the ISCO of \\j\\ in other works) throughout the transition from the \\textit{brighter} phases of the low-hard state to the intermediate states (both the hard-intermediate and soft-intermediate), through to the soft state and back; (2) the ratio between the observed reflected X-ray flux and power-law continuum (the ``reflection fraction'', $R$) increases sharply at the transition between the hard-intermediate and soft-intermediate states (``ballistic'' jets are sometimes launched at this transition); (3) both the frequency and coherence of the high-frequency quasi-periodic oscillations (QPOs) observed in \\j\\ increase with $R$.  We discuss our results in terms of black hole states and the nature of black hole accretion flows across the mass scale. ",
        "introduction": "\\label{introduction}   The majority of stellar-mass black holes residing in low mass binary systems spend most of their lives in quiescence. This explains the empirical fact that out of tens of thousands of such systems predicted to exist throughout our Galaxy \\citep[e.g.][]{Yungelson2006}, only about 50 have been discovered. For a number of such systems, it is well established that the outburst -- which resulted in their discovery -- evolves through a number of spectral states characterised by the relative strength of their thermal and non-thermal X-ray emission and with possible differences in the accretion  geometry and reflection attributes. These active states can be roughly separated into four semi-distinct states which are phenomenologically described below and extensively discussed in \\citet{Remillard06} and \\citet{Bellonibook2010}. At the outset of the outburst, the system goes through what has been traditionally dubbed the low-hard state (LHS) where the X-ray spectrum is dominated by a non--thermal component often simply described by a power-law (photon index $\\Gamma$ between $\\sim1.4-2$)  spectrum with relatively low luminosity ($\\sim 0.05 L_{\\rm Edd} $ and and exponential cut-off at $\\sim100$\\kev).   The  energy spectrum in the LHS peaks near $\\sim 100$\\kev\\ and often we also see the  weak presence of a thermal component (contributing  $<20\\%$ of the total 2-20\\kev flux)  with a temperature below $\\sim0.5$\\kev\\  produced by the accretion disk \\citep[see e.g.][and references therein]{Reynold2011swift, reislhs}.   As the luminosity increases, the spectrum moves through the intermediate state (IS)  where the 2-10\\kev\\ flux    is  typically a factor of $\\sim4$ times higher than that of the LHS. Here, the soft ($\\Gamma = 2-3$) power-law tail coexists with a strong thermal component. Recently, the intermediate state has begun to be subdivided into an early Hard-Intermediate State (HIS) and a later Soft-Intermediate state (SIS)  just prior to a transition into the canonical High-Soft or thermal state, where  the X-ray flux is dominated ($>75\\% $ of the total 2-20\\kev flux) by the thermal radiation from the inner accretion disk having an effective temperature of  $\\sim1$\\kev. In this final state of the outburst, the system usually emits with luminosities $>0.1 L_{\\rm Edd}$ and the  power-law component is both weak (less than 25\\% of the total 2-20\\kev\\ flux) and steep ($\\Gamma = 2-3$).  Following the HSS, the system often returns to the LHS and subsequently goes back to quiescence were it can remain indefinitely or in some cases for a few years before this cycle restarts.    The hard X-ray emission predominant in the LHS has long been linked to inverse Compton scattering of the soft thermal disk photons by a population of hot  ($\\sim10^9\\k$) electrons in a cloud of optically thin, ionised gas or ``corona\" surrounding the inner parts of the accretion disk \\citep{Shapiro1976,SunyaevTitarchuk1980}.  Under the common assumption that the radio emission observed to originate from stellar mass black holes, is directly related to the presence of a jet, it is believed that all such systems, either in the LHS or in transition, launch a collimated  outflow \\citep[e.g.][]{Fender2001jets, fenderetal04, Fender091}. The fact that these persistent jets are observed only in the LHS suggests that the jet is linked to the corona, with claims that the corona in the LHS is indeed the launching point of persistent  jets \\citep[see e.g.][]{Markoff05}. The connection between the radio (jet) and X-ray flux for both stellar-mass and supermassive black holes \\citep{fundamentalplane,GalloFenderPooley2003,fundamentalplane2}, often referred to as ``the fundamental plane\" of black hole accretion, suggests an intimate connection between the corona and radio-jets \\citep[see e.g.][]{miller2012cyg}.  Whether state transitions are driven by intrinsic changes in $\\dot{m}$, physical changes in the disk,  disk-corona, radio jet or a combination of all these factors is a matter of much debate.    \\subsection{Reprocessed X-rays: Reflection } The existence of a hard X-ray source -- the corona -- also adds further complexities to the various spectral states. The reprocessing of these hard X-rays by the relatively cold accretion disk in all active states results in a number of ``reflection features\" consisting of discrete atomic features together with a ``Compton-hump\" peaking at approximately 30\\kev. The high fluorescent yield -- and relatively high cosmic abundance -- of iron often results  in a particularly strong feature at $\\sim 6-7$\\kev\\ \\citep[see e.g.][for a recent review of  ``reflection\" in black holes]{Fabianross2010}.   The strong irradiation of the black hole accretion disk by the coronal photons likely causes the surface layers to be photoionised. \\citet{rossfabian1993} investigated the effect of allowing the gas constituting the top layers of the accretion disk to ionise, and the authors went on to compute reflection spectra for different ionization levels. A number of similar studies of reflection from ionised matter have been conducted since \\citep{MattFabianRoss1993, MattFabianRoss1996, rossfabianyoung99,NayakshinKazanasKallman2000, NayakshinKallman2001, cdid2001, GarciaKallman2010, GarciaKallman2011}. These studies demonstrate that  the reflection spectrum expected from a black hole depends strongly on the level of ionization of the surface layers of the disk. This can be quantified for a constant density gas by the ionization parameter \\begin{equation}\\xi  = \\frac{L_{\\rm x}}{nd^2}, \\end{equation} where $L_{\\rm x}$ is the ionising luminosity of the source, $d$ is the distance between the disk and the source, and $n$ is the density of the disk. Thus an increase in $\\xi$, either by increasing the illuminating flux, decreasing the density or distance between the X-ray source and the disk, will cause the gas in the disk to become more ionised. \\citet{MattFabianRoss1993, MattFabianRoss1996} split the behaviour of the reflection spectrum into four main regimes depending on the value of $\\xi$.   For low ionization parameter ($\\xi < 100~\\ergcmps$),   the material is only weakly ionised and the reflection spectrum resembles that arising from  ``cold'' matter, with a prominent iron line at 6.4\\kev, and strong absorption below $\\approx 10$\\kev. There is only a weak contribution from the backscattered continuum at $\\approx 6$\\kev\\ and a weak iron K absorption edge at 7.1\\kev. As the disk becomes more ionised ($100 < \\xi < 500~\\ergcmps$) the system reaches  intermediate ionization range where Fe has lost all of its M-shell ($n=3$) electrons and thus exists in the form of FeXVII--FeXXIII with a vacancy in the L-shell of the ion. Due to this vacancy, the L-shell can absorb the \\ka\\ line photons and thus effectively trap the escaping photon. This resonance trapping is only terminated when an Auger electron is emitted. This second  ionization regime is therefore characterised by a very weak iron line and a moderate iron absorption edge.  As the gas becomes more ionised ($500 < \\xi < 5000~\\ergcmps$) all low-$Z$ metals are found in their hydrogenic form and  the soft reflection spectrum has only weak spectral features. Iron is found mostly in its hydrogen or helium like forms (FeXXVI or FeXXV respectively) and due to the lack of at least 2 electrons in the L-shell (i.e. not having a full 2s sub-shell), Auger de-excitation cannot occur. The result is strong emission from ``hot'' \\ka\\  FeXXV and FeXXVI at 6.67 and 6.97\\kev\\ respectively and the corresponding absorption edges at approximately 8.85 and 9.28\\kev\\ respectively. Finally, when $\\xi \\gg 5000~\\ergcmps$, the disk is  highly ionised and there is a distinct absence of any atomic features.    A further complication arises in the reflection spectra of stellar mass black holes due to the fact that in these systems the gas in the accretion disk is inherently X-ray ``hot'' meaning that low-$Z$ metals can be fully ionised in the gas even  before receiving any irradiation by the X-ray corona. To account for this extra ``thermal ionization'', \\citet{refbhb} performed self-consistent calculations of the reflection resulting from the illumination of the accretion disk by both a hard, powerlaw corona and thermal disk blackbody radiation. The authors compared the results of having the disk both in hydrostatic equilibrium and under the assumption of a constant density atmosphere, and found reasonably  good agreement between the two emergent spectra. \\citet{refbhb} also confirmed in stellar mass black holes the result that had been previously found for AGN in that the spectrum from a constant density disk is slightly diluted (it has a lower  flux)  in comparison to that of a disk in hydrostatic equilibrium. Furthermore, the authors also noted a few small differences between the modes; namely a lower  effective ionization parameter in  the constant density model which resulted in a slightly stronger Fe\\ka\\ line and deeper iron K-edge. Nonetheless, the overall spectrum from the constant density approximation was shown to be in  good agreement with the result for an atmosphere in hydrostatic equilibrium. The two reflection grids resulting from the work of \\citet{reflionx, refbhb}, will be used frequently throughout this work.   \\subsection{General Relativistic Effects: Light bending}  Naively,  assuming  isotropic coronal emission, one would expect variations in the  reflection component to be directly correlated to variations in the observed power-law continuum.   That is, as the observed flux of the  X-ray corona increases, so should the amount of reprocessed emission. However, in a number of instances it has been found that this is not the case, with the reflection component at times behaving in an anticorrelated manner \\citep[e.g.][]{Rossi2005j1650} or not varying at all despite large variations in the X-ray powerlaw continuum \\citep[e.g.][]{Fabian02MCG, Fabianvaughan03, MiniuttiFabianGoyder2003,BallantyneVaughan2003,LarssonFabian2007}.   By virtue of its proximity to the black hole, the emission from the corona is naturally  affected by general relativistic (GR) effects. Some of the radiation from the corona which would otherwise escape is gravitationally focused  -- ``bent\" --  towards  the accretion disk giving rise to enhanced reflection and selectively decreasing the X-ray continuum at infinity. A number of studies \\citep[for instance][]{MartocchiaMatt1996, MartocchiaKaras2000LB, Miniu04, MiniuttiFabianMiller2004j1650, Niedzwiecki2008} have investigated  the effect of GR on a compact,  centrally concentrated X-ray corona close to a black hole\\footnote{Observational  evidence for such compact X-ray corona has recently come from microlensing results where the size of the X-emitting region has been shown to be of the order of $\\sim10$\\rg\\ \\citep{ChartasKochanek2009quasar,DaiKochanek2010quasar, ChartasKochanek2012quasar, MorganHainline2012quasar}.}. The ``light-bending\" model put forward by  \\citet{Miniu04} predicts a number of semi-distinct regimes affecting the variability of the reflection component compared to the X-ray continuum:  \\begin{description} \\item[Regime 1]  When the corona is very close to the black hole (a few gravitational radii \\rg=$GM/c^2$), a large fraction of the radiation is bent onto the the accretion disk thus significantly reducing the amount of observed X-ray continuum and enhancing the reflection. A very steep emissivity profile is expected as the source is highly concentrated in the inner regions and the reflection is expected to be a steep function of the continuum in a quasi-linear manner.  \\item[Regime 2] When the central corona is slightly further from the black hole (at heights of $\\sim 10$\\rg), light bending causes the reflection component to vary significantly less than the  X-ray continuum. The amount of light bent towards the black hole decreases as the corona moves further from the black hole and the X-ray continuum increases.   \\end{description}    Finally, at heights  $\\gg 20$\\rg, light-bending becomes less important and the observed  continuum increases\\footnote{ In the original paper of \\citet{Miniu04}  a further regime -- Regime 3 -- was defined at large radii where the reflection and powerlaw flux appeared to be anti-correlated with one another. This was an artefact of  having a finite boundary for the disk extent of 100\\rg, and instead the reflected flux should asymptotically become flat with respect to the continuum as a continuation of Regime~2 \\citep[e.g.][]{Niedzwiecki2008}.}.    In this manner, the presence of gravitational light-bending has been invoked to explain the fact that Seyferts  \\citep[e.g.][]{FabZog09, Fabian20121h0707} and XRBs \\citep[e.g.][]{reismaxi}  at times appears to be ``reflection-dominated\". Sources where the observed X-ray spectrum display a distinct lack (or comparably small amount) of hard, powerlaw-like continuum despite a large contribution of reflection are fully consistent with the first Regime detailed above.   Of course, the model presented above is idealised in that all characteristics of the observed variabilities are assumed to be a result of variation in the height of an isotropic  and compact corona with a fixed luminosity. Although intrinsic variation in the luminosity of the corona may well be present, it is unlikely that they could solely explain the behavior of the reflected emission described above. Indeed, the  clear presence of broad and skewed iron  emission lines in a growing number of sources ranging from stellar mass black holes \\citep{miller07review,miller09spin, reisspin, Steiner2011, hiemstra1652} to AGNs \\citep{tanaka1995,Nandra97, Nandra07, FabZog09,3783p1} (and also to lesser extent neutron stars \\citep{cackett08, cackett10, DiSalvo2005170544, disalvo09, reisns}) strongly attest that general relativistic effects play an important role in producing the line profile, further  supporting the notion that the corona and the inner disk are both in the inner regions around a black hole.       \\subsection{\\j}   Amongst one of the first systems to provide observational evidence for the aforementioned  effect of gravitational light bending around a black hole and indeed the first around a stellar mass black hole was \\j, which was discovered  by \\rxte\\ on 2001 September 5 as it went into outburst  \\citep{Remillard2001}. Based on the spectrum obtained early in decay of the outburst by \\xmm, and more importantly on the presence of a clearly broad iron emission line, \\citet{miller02j1650} were able to not only infer that the object in \\j\\ was indeed a black hole, but also that it was close to maximally rotating with a dimensionless spin parameter $a$ $\\approx0.998$.  By decomposing the hard X-ray continuum from the reflection component in three \\bepposax\\ observations of \\j, \\citet{MiniuttiFabianMiller2004j1650}  were able to show that the  latter remained nearly constant despite a large change in the direct continuum, in a manner consistent with the predictions of light bending around a black hole.  Optical observations obtained after the system had returned to near quiescence  \\citep{Orosz2004J1650} revealed a mass function $f(M) = 2.73\\pm 0.56\\msun$ with a most likely mass of $\\sim4\\msun$, and with it secured \\j\\ as genuine black hole  binary system. \\citet{CorbelFender2004} reported on the radio and X-ray observations of \\j\\ during the outburst. The authors find a clear drop by nearly an order of magnitude in the radio flux at the transition from the hard intermediate state (referred to as the intermediate state in that work)  to the soft intermediate state (referred to as the steep power-law state in that work), and surprisingly they find residual radio emission during the often radio quiet disk-dominated soft state  which they attributed  to possible emission from previously ejected material interacting with the interstellar medium, rather than originating in the central source.   A follow up study of \\rxte\\ data by \\citet{Rossi2005j1650} used the iron line as a proxy to the total reflection component and confirmed the plausibility of the light-bending scenario for the evolution of \\j. Again using data obtained from the \\rxte, \\citet{Homan2003} reported on the discovery of a $\\sim250$Hz QPO together with a number of less coherent variability peaks at lower frequencies. By studying the spectral and timing evolution during the first $\\sim 80$ days of the outburst, the authors were able to define six  periods (I--VI; in this work referred to as P1--P6) having somewhat distinct   spectral and timing characteristics (see their Fig~1 and Table~1). A recent study involving  \\j\\  has discussed the similarities between X-ray binaries and AGNs \\citep{Waltonreis2012} and argued that both \\j\\ and the active galaxy MCG--6-30-15 \\citep{tanaka1995} must contain a rapidly rotating black hole, with the spin of \\j\\ having been formally constrained to $0.84\\leq a \\leq 0.98$.  A further body of work based on \\rxte\\ observations and a variety of empirical models for the hard X-ray continuum   \\citep{YanWang2012} has concluded that the emission region, here referred  to as the corona, decreases by a factor of $\\sim 23$ in size during the transition from the hard to the soft state.  \\subsection{This Work}  \\begin{figure*}[!t] \\centering { \\rotatebox{0}{ {\\includegraphics[height=5.5cm]{fig_lc_bkn1-eps-converted-to.pdf}}}} \\hspace{-0.3cm} { \\rotatebox{0}{ {\\includegraphics[height=5.5cm]{fig_lc_bkn2-eps-converted-to.pdf}}}}  \\vspace{-0.2cm} \\caption{\\label{fig1} The 181 observations used were taken during the 2001/2002 outburst which is clearly seen in the All Sky Monitor (left). Since its discovery outburst, \\j\\ has remained in quiescence. (Right:) PCA-PCU2 (top) and HEXTE-A (bottom) count rate during the time encompassing the outburst. The different colors marks the distinct spectral states as defined by Homan et al.~(2003) based on the timing characteristics of the source. }  \\vspace*{0.2cm} \\end{figure*}  Since its discovery, \\j\\ has become one of the best studied black hole systems. However, the energy spectra of this system have either been studied in a great degree of detail using high-quality, single snapshot observations with \\xmm\\ \\citep[i.e.][]{miller02j1650, Waltonreis2012} or \\bepposax\\ \\citep{MiniuttiFabianMiller2004j1650} or using mostly phenomenological and simple models in the study of long term evolutions with \\rxte\\ \\citep[i.e.][]{Rossi2005j1650, Dunn2010, Dunn2010disk, YanWang2012}.  In this paper, we use the full \\rxte\\ archival data of the outburst to {investigate the evolution of the direct power-law continuum, reflection and thermal disk  components using, for the first time,  a fully self-consistent prescription for the reflection component}.  In this manner, we combine the virtues of detailed analyses of single observations with the immense diagnostic power of multiple \\rxte\\ pointings. By being able to  decouple the total reflection component (Fe-\\ka\\ emission line together with all other reflection signatures) from the illuminating continuum, we find that the transition from the hard-intermediate state to the soft-intermediate state is accompanied by a sharp increase in the strength of the reflected emission in comparison to the direct continuum. We interpret this increase in the reflection fraction as a sudden collapse of the corona as the system approaches the thermal state, although we note that this may not be a unique interpretation.   This paper is structured as follows: \\S2 briefly introduces the observations and details our various selection criteria. \\S3 describes the base model and assumptions used throughout this work. The various results are presented in \\S4 and in \\S5 we present a qualitative picture of a possible physical scenario that combines all our findings. ",
        "conclusions": "\\subsection{Evolution of the Outburst} \\label{Light bending}   \\begin{figure*}[] \\centering  \\hspace*{0.cm} {\\hspace{-0.0cm}\\rotatebox{0}{{\\includegraphics[width=5.5cm]{figures_spec1-eps-converted-to.pdf} }}} {\\hspace{0.cm} \\rotatebox{0}{{\\includegraphics[width= 5.5cm]{figures_spec6-eps-converted-to.pdf} }}} {\\hspace*{0.cm} \\rotatebox{0}{{\\includegraphics[width= 5.5cm]{figures_spec22-eps-converted-to.pdf} }}} {\\hspace{-0.0cm}\\rotatebox{0}{{\\includegraphics[width=5.5cm]{figures_spec46-eps-converted-to.pdf} }}} {\\hspace{0.cm} \\rotatebox{0}{{\\includegraphics[width= 5.5cm]{figures_spec64-eps-converted-to.pdf} }}} {\\hspace{0.cm} \\rotatebox{0}{{\\includegraphics[width= 5.5cm]{figures_spec86-eps-converted-to.pdf} }}} {\\hspace*{0.cm} \\rotatebox{0}{{\\includegraphics[width= 5.5cm]{figures_spec101-eps-converted-to.pdf} }}} {\\hspace{-0.0cm}\\rotatebox{0}{{\\includegraphics[width=5.5cm]{figures_spec102-eps-converted-to.pdf} }}}  \\caption{Unfolded spectra (top) and residuals (bottom) for the eight representative states shown in Figs.~1,2.  The reflection, power-law, and disk component are shown in red, blue and green respectively. The total model is shown in black.  All vertical scales are the same except for that of the falling LHS (Bottom right). The approximate MJD of each observation are shown in brackets. }  \\vspace*{0.3cm} \\label{fig4} \\end{figure*}     Figure~3 shows the evolution of all the parameters of interest in this work\\footnote{We refer the reader to the work of \\citet{Dunn2010disk} for the evolution of the disk properties during the outburst. As mentioned above, despite the fact that the authors did not correctly account for reflection in their work, this plays a minor role at the energies of interest in regards to the disk properties and as such that work is still a valid and important reference for the evolution of accretion disks.}, and Fig.~4 shows the best-fit to eight representative spectra roughly covering the eight periods highlighted in Figs.~1,2 and described in detail in Homan et al.~(2003). The spectra used for illustration are shown in Fig.~2 with diamonds.  The top four panels in Fig.~3 show the evolution of  the extrapolated 0.1-1000\\kev\\ fluxes for the total, \\po, \\reflionx\\ and \\diskbb\\ components, from top to bottom respectively. All fluxes were obtained using  \\cflux\\ in \\xspec.  The vertical dotted lines running through all the panels highlight the eight periods shown in Figs.~1,2. We see a clear increase in the disk flux during the first $\\sim30$ days followed by a clear flattening as it moves into the HSS. It is also visually apparent that the reflection flux varies relatively less than the power-law continuum. This will be investigated further in what follows.  The following  two panels show the evolution of the disk temperature and  ionization parameter, $\\xi$. Early in the outburst, the disk was relatively cold ($\\lesssim0.5$\\kev) and only moderately ionised with $\\xi\\approx200\\ergcmps$. As the system moves into the HSS, the ionization increases smoothly until  $\\sim2$~days before the transition to the SIS when $\\xi$ sharply increases to $\\xi\\approx3000\\ergcmps$ and remains at that level through the transition up to the end of the SIS. The disk temperature, on the other hand, appears to reach a relativity  stable value of $\\sim0.6-0.65$\\kev\\ early in P2, approximately half way through the HIS. As the system progress into the HSS, the disks becomes  fully ionised with $\\xi$ reaching the maximum allowed value in the model (log~$\\xi=4$), before coming back down to the low hundreds towards the end of the outburst.  The reflection fraction $R$ shown in the third-from-bottom panel of Fig.~3 is here defined as a measure of the ratio between the (observed) continuum power-law  to the reflection flux emitted by the disk. Since a fraction of the  power-law illuminating the accretion disk is down-scattered as it is reprocessed in the disk, the reflection fraction in \\reflionx\\ is calculated by dividing the  extrapolated  (1\\ev - 1000\\kev) \\reflionx\\ flux by the   0.1-1000\\kev\\ power-law flux.  At the start of the outburst, through to the end of the HIS, $R$ increases smoothly between   $\\approx 0.6-1$. At the transition between the HIS to the SIS, $R$ displays a sharp increase to $\\approx  4$ where it remains until the beginning of the disk dominated HSS where the power-law effectively disappears. We also show in this panel the ratio of the 3-100\\kev\\ reflection to powerlaw flux. The behavior described above is still qualitatively the same and we still see a clear jump in this ratio at the transition from the HIS to the SIS. However, when limited to the 3-100\\kev\\ flux, this ratio is systematically a factor of $\\approx1.5$ less than the extrapolated ratio;  a direct  result of not accounting for the extra down-scattered flux at low energies.  As a further test of both the qualitative (clear jump in \\textit{R} between HIS and SIS) as well as quantitative (change from $R\\approx 0.6$ to $R\\approx 4$ between P1 and P3)  behaviors found here for the reflection fraction, we temporally replace the \\reflionx\\ model with a combination of \\laor\\ plus \\pexriv\\ and employ this model to the observations highlighted in Figs. 2 and 4 for P1 and P3. In using this model, we have blurred the \\pexriv\\ component with the same parameters as the \\laor\\ line profile Figure~5 shows the confidence range for the reflection fraction (a free parameter of \\pexriv)  for these two representative spectra. In agreement with our previous results, we see that in the HIS the reflection fraction is constrained to  $R=0.58^{+0.08}_{-0.11}$ and in the SIS it is $R>2.7$ at the 90~per~cent level of confidence ($\\Delta\\chi^2 = 2.71$).  \\begin{figure}[] \\centering \\vspace{0cm} {\\hspace*{-0.2cm}\\rotatebox{0}{{\\includegraphics[width=9cm]{figures_pexriv-eps-converted-to.pdf} }}}  \\caption{Goodness-of-fit versus reflection fraction for the two representative spectra describing the HIS-P1 (black) and SIS-P3 (red). The spectra used refers to those highlighted in Figs. 2 and 4 and the \\reflionx\\ component inherent in the base model has been replaced with a combination of \\pexriv\\ together with \\laor. The solid blue horizontal line shows the 90~per~cent confidence range. } \\label{fig5} \\end{figure}   \\begin{figure*}[] \\vspace*{0.0cm} \\centering {\\hspace{-0.0cm} \\rotatebox{0}{{\\includegraphics[width=14cm]{figure_reducedchi_histogram4-eps-converted-to.pdf} }}} \\vspace*{-0.1cm}  \\caption{Left: Distribution of reduced $\\chi^2$ for model with a Newtonian emissivity profile ($q=3)$ and a steeper values $\\geq 3$. In both cases we see a peak at reduced $\\chi^2 =1$, however this is much clearer after allowing the emissivity to vary beyond its Newtonian value. Right: Distribution of emissivity parameter (blue) and their 90\\% lower limit (red) for the four states highlighted in each panel Bottom:  In all cases, the emissivity index was constrained to $3\\leq q \\leq 10$.   }  \\vspace*{0.2cm} \\label{fig6} \\end{figure*}   The fact that $\\xi$ is maximal in the HSS (P4) despite the fact that this is when the  irradiation appears to be at its lowest level (second panel from the top) can be explained with a number of scenarios. We showed in \\S~1.1 that in stellar mass black holes the intrinsic hot disk can result in significant thermal ionization \\citep{refbhb} which will be strongest in the disk dominated HSS. Thus, in this scenario, the high ionization measured could also be in part due to thermal ionization.  If this thermal component is the sole source of ionisation in the HSS, $R$ would indeed go to zero.   A further possibility is that the disk is indeed highly photo-ionised as a result of strong focusing of the coronal photons onto the disk. This  would significantly remove the number of hard-photons escaping the system (thus explaining  the second panel from the  top) and cause the disk to be highly ionised. Indeed, observations of disk winds originating in the HSS of various BHBs consistently show winds having  $\\xi \\sim 10^4 \\ergcmps$ \\citep[e.g.][]{Miller2008j1655wind,NeilsenLee2009,KingMiller2012,Ponti2012diskjet}. Unfortunately, the lack of reliable constrains on the reflection fraction in the HSS prevents us from making any solid claims on the nature of the disk-corona interaction in this state.  \\subsection{Disk Emissivity}      In Fig.~6  (left) we show the distribution  of the reduced $\\chisq$ (for 115 degrees of freedom) from all observations assuming the simple Newtonian  `lamp-post'  like geometry in which the  the emissivity profile follows a $q=3$ power-law (light blue),  as well as after relaxing this assumption (red). In both cases there is a clear peak at reduced-$\\chisq = 1$, however this peak is much more distinct upon relaxing the Newtonian approximation. The Newtonian approximation naturally does not take into account the effects of general relativity that will  be experienced by the emission from the corona and accretion disk near the black hole. Relativistic effects (strong gravity as well as relativistic time dilation) acts to steepen the emissivity in the inner regions of the disk.  The right panels of Fig.~6 show either the distribution of the emissivity index (blue) or the 90~per~cent lower limit  in their value (red) for the spectral states indicated in each panel.  We refer the reader to the work of \\citet{Miniu04, wilkins2011, wilkins2012, fabian2012cyg} and references therein for a detailed study examining non-Newtonian values for the emissivity index, but note here that steep emissivities similar to those found here for the HIS and SIS are a natural and unavoidable consequence of strong gravity.    The bottom panel of Fig.~6 is used here as a simple illustration of the evolution in $q$ as well as the count-rate in both the PCA and HEXTE data. It is clear that $q$ can only be constrained when the PCA data is at its highest level as this constraint does not come from energies $>25$\\kev. At the end of the outburst, when the PCA signal-to-noise level drops significantly, the data cannot differentiate between a Newtonian   $q=3$ and a  steeper value.   Following the recipe provided for Cygnus X-1 by \\citet{fabian2012cyg} in dealing  with sources where the spin is expected to be high (as is the case for \\j), we have repeated our fits with a double emissivity profile such that within a break radius (initially frozen at 4\\rg\\ but later allowed to vary) the emissivity is $>3$ and beyond it is frozen at 3. The initial value of 4\\rg\\ for the break radius was chosen based upon the value for Cygnus X-1  \\citep{fabian2012cyg}.  We find that, as long as the emissivity is not fixed at  $q=3$, the quality of the fits, and distribution of reduced $\\chisq$ are similar to that of a single, unbroken emissivity, and we proceed by using this single power-law emissivity profile as our standard but emphasise that the results presented here do not change if we employ a broken emissivity profile. This is very likely to be due to the comparatively low spectral resolution afforded by RXTE which does not reflect the subtle changes in the reflection profile in a similar manner as \\xmm\\ or \\suzaku\\ observations.     \\begin{figure*}[!t] \\vspace*{-0cm} \\centering {\\hspace{-0.0cm} \\rotatebox{0}{{\\includegraphics[width=14.6cm]{fig_flux_flux3-eps-converted-to.pdf} }}} \\vspace*{-0.0cm} \\caption{ Left: Flux-Flux relation during the outburst. The dashed line shows the one-to-one relation. Top-right: Close up of the HIS (blue and cyan) and SIS (green).  The solid black lines shows the best linear fit for each state having a slope of $2.7\\pm0.1$ and $-0.24\\pm0.02$ in the SIS and HIS respectively. The solid red curve shows the expected flux-flux relation under the light bending model of \\citet[][see their Figure 2]{Miniu04} for a system with 60\\deg\\ inclination, somewhat similar to \\j, with a corona varying in height from 1--20\\rg. Note that this is not a fit and has been rescaled from the original (see \\S~4.3). Bottom-right: Similar but for fluxes between 3--100\\kev.  } \\label{fig7} \\end{figure*}    \\begin{table*} \\vspace{-0.4cm} \\caption{Summary of  Spearman's Rank Correlations and Partial Correlation tests on a combination of various model parameters. } \\centering \\begin{tabular}{cc|cccc|cccc} \\hline \\hline \u2022 & \u2022 & \\multicolumn{4}{c|}{HIS (P1 + P2)} & \\multicolumn{4}{c}{SIS (P3)}\\\\ \\hline \u2022 & \u2022 & \\multicolumn{2}{c|}{Spearman's}   & \\multicolumn{2}{c|}{Partial } & \\multicolumn{2}{c|}{Spearman's}   & \\multicolumn{2}{c}{Partial } \\\\ \u2022 & \u2022 & \\multicolumn{2}{c|}{rank-order}   & \\multicolumn{2}{c|}{correlation } & \\multicolumn{2}{c|}{rank-order}   & \\multicolumn{2}{c}{correlation } \\\\ \\hline Parameter 1 & Parameter 2 & $\\rho$ &  p-value & $\\rho$ & p-value  & $\\rho$ & p-value  & $\\rho$ & p-value  \\\\ \\hline  $\\xi$ & $F_{\\rm powerlaw}$ & -0.953 & $ 2.0\\times 10^{-13}$ & -0.473 &  0.014 & 0.319 & 0.071 & 0.092& 0.621 \\\\  $\\xi$ & $F_{\\rm disk}$        & 0.956 & $ 9.3\\times 10^{-14}$ & 0.515 &0.006 & 0.006 & 0.972 & 0.082& 0.656 \\\\  $F_{\\rm disk}$ & $F_{\\rm powerlaw}$ & -0.955 &$ 1.1\\times 10^{-13}$ &-0.427 & 0.0031 & -0.310 & 0.079 &  -0.334 & 0.056 \\\\  $\\xi$ & $F_{\\rm reflionX}$ & 0.742 & $ 2.2\\times 10^{-5}$ & -0.140 & 0.516 & 0.382 & 0.028 & 0.218 & 0.230\\\\  $F_{\\rm disk}$ & $F_{\\rm reflionX}$ & 0.766 & $ 7.8\\times 10^{-6}$ & 0.088 & 0.685 & -0.116 &0.520  &0.140 &0.446 \\\\  $F_{\\rm powerlaw}$ & $F_{\\rm reflionX}$ & -0.799 & $ 1.7\\times 10^{-6}$ & -0.373 & 0.066 & 0.719 & $2.4\\times 10^{-6}$ & 0.683 &  $ 4.6\\times 10^{-7}$ \\\\ \\hline \\end{tabular}  \\vspace{0.2cm} Notes:-  Spearman's rank correlations and Partial Correlation test were made for combinations of the reflection, power-law and disk fluxes as well as the ionization parameter. The Partial Correlation test measures the degree of associations between the two parameters listed on the first two columns whilst controlling for the remaining two parameters. The Spearman's coefficient $\\rho$ is a measure of the degree of correlation with +1 or -1 indicating a perfect monotone function and 0 a lack of correlation.\\vspace*{0.2cm} \\end{table*}      \\subsection{Light-Bending and General Relativity}  Hints of  the expected effects of light bending, as described in the introduction, can be seen in the top four panels in Fig.~3.  Most important, is the apparent constancy of the reflection flux in comparison with that of the direct power-law early in the outburst. We investigate this further in Fig.~7. The left panel shows the flux-flux relation  through the whole outburst with the various spectral states shown in different colors. The top-right panel is a close up of the period covering the  HIS and SIS during the first $\\sim$~30 days of the outburst\\footnote{Excluding the first 3 days when \\j\\ was in the rising LHS (see Fig.~1).}. Figure~7 is remarkably similar to figure~3 of  \\citet{Rossi2005j1650}, where the authors used the flux in the iron line as a proxy for the total reflection in \\j. We have superimposed in this figure the expectation from the light bending model for a compact corona varying in height from 1--20\\rg\\ with a disk having an inclination of 60\\deg, from \\citet[][see their Figure 2]{Miniu04}. In order to correctly describe the shape of the function shown graphically in \\citet{Miniu04}, we used the  Dexter Java application of \\citet{dexter_ads} to obtain a fourth-order polynomial fit to their curve from which we applied a linear normalisation of $1.5\\times10^{-9}$ and $4.5\\times10^{-9}$ to their Y (arbitrary Fe line flux) and X-axis (arbitrary powerlaw flux) respectively.  The model reproduces extremely well the broad  shape of the relation. Finally, the bottom right panel shows this behavior when the non-extrapolated, 3-100\\kev\\ fluxes are used instead\\footnote{In this case the normalisations were $1.5\\times10^{-9}$ and $4.5\\times10^{-9}$ for the Fe-line and powerlaw flux, respectively). }. We again see that qualitatively, the behavior is the same as above.   As discussed in the introduction, the light bending model of Miniutti et al.~(2004) predicts the existence of semi-distinct regimes in this flux-flux relation. When the corona is located at a height of  $\\sim 10$\\rg\\ the model predicts a flattening in the relation similar to that observed for the HIS (both P1 and P2). The fluxes in this hard intermediate state are clearly correlated, and a Spearman's rank correlation test gives a coefficient of $\\rho = -0.799$ corresponding to  a $1.7\\times10^{-6}$ chance of a false correlation (Table~1).  The slope of this relation is $-0.24\\pm0.02$ (standard error, s.d.)  and this linear fit is shown in the top-right panel of Fig.~7 as a solid black  line.  As the location of the corona reaches the more extreme environment  within a few \\rg\\ from the black hole, the model predicts a steep, positive linear relation between the reflection and power-law flux similar to that seen in the SIS, although there is large scatter dominated by poor statistics.  A Spearman's rank correlation test here gives $\\rho = 0.719$, with a false correlation probability of just $2.4\\times10^{-6}$. The slope of this relationship is  greater than unity, with a best-fit value of $2.7\\pm0.1$ (shown as a further solid black  line in Fig.~7) . Note that this is highly inconsistent with the expectation for a static Newtonian corona  with intrinsically varying luminosity, where the slope should be unity across the entire flux range.   It is the combination of a slope~$\\gg1$  at low powerlaw flux together with a  near-flattening at higher fluxes that provides evidence for relativistic effects in this case.   It has been suggested \\citep[e.g.][]{BallantyneVaughan2003, Ballantyne2011} that changes in the ionization of the inner regions of the accretion disk can give rise to changes the reflection spectrum that can mimic somewhat this flat behavior. In order to robustly assess the strength of the correlations seen in Fig.~7 in both the HIS and SIS, we have performed a number of correlation tests which are summarised in Table. 1.  We performed, for both HIS and SIS,  Spearman rank-order tests for a combination of four parameters (power-law, reflection and disk fluxes as well as the disk ionization) as well as Partial Correlation tests (PCT) for two parameters while controlling for the third and fourth variable. The PCT is of particular importance for our purpose as it removes any potential association  of the ionization parameter (or any other potential source of unwanted correlation) in the flux-flux relations shown in Fig.~7.  From the  Spearman rank-order tests performed in the HIS (Table~1), it would initially appear  that all four variables are strongly correlated with one another in some way, as all combinations display $|\\rho|\\gtrsim0.7$. However, after performing the partial correlation test for all combinations we see that for  two of the  previous strong correlations ($\\xi-F_{\\rm reflionX}$ and $F_{\\rm disk}-F_{\\rm reflionX}$)  were in fact driven by the  mutual dependence of these parameters on $F_{\\rm powelaw}$. These tests clearly indicate that the reflection flux in both states is better correlated with the power-law flux than with ionization and the slope of the correlations (mildly negative in the HIS and strongly positive in the SIS) are highly indicative of  gravitational light-bending in the General Relativistic~regime.   The behavior seen here is consistent with a drop in the height of the corona during the hard-intermediate phase (P1 and P2) followed by intrinsic variations in its luminosity by a factor of a few during the soft-intermediate state (P3). Following the disk-dominated soft-state (P4), the height of the corona increases again (P5 and P6) and the outburst finishes with the intrinsic power dropping as the source goes back into the LHS.    \\subsection{$R-\\Gamma$ relation}  A strong correlation has been shown to exist between the amplitude of the reflected component ($R$) and the photon index ($\\Gamma$)  of the Comptonized spectrum in XRBs {\\it in the hard state}  \\citep[e.g.][]{Ueda1994rgamma}. This $R-\\Gamma$ relation has since been robustly tested by a number of authors \\citep[e.g.][]{Gilfanov1999rgamma, Zdziarski1999, Zdziarsk2003, NowakWilmsDove2002gx, ibragimov05} and it is now thought to also apply to Seyferts and radio galaxies \\citep[e.g.][]{Zdziarski1999},  further cementing the similarities in the coronal properties at all scales. If this relation indeed turns out to be real (and evidence attests  to this; but see \\citealt{Molina2009rgamma}) then this is could be telling us about the  feedback between the hot corona and cold gas in an accretion disk.  We briefly investigate this relationship for \\j\\ in Fig.~8. For clarity, we only consider data with a fractional uncertainty of less than 50~per~cent. The black arrows approximately show the evolution of the system in time. Although as a whole the data does not strongly support the presence of a relationship between $R$ and $\\Gamma$, when the states are roughly separated (different colors) it does appear that early in the outburst through to the  last few days of the hard-intermediate state the relation seems to hold. It is clear, at least, that the rising-LHS and the HIS populate different regions in the figure to the SIS.  The potential presence of  this relation early in the outburst  suggests a  feedback process between the soft photons in the disk and the corona. There are a number of theoretical interpretations for the presence of this correlation \\citep{PoutanenKrolik1997,Gilfanov1999rgamma,Gilfanov2000rgamma, beloborodov99, MalzacBeloborodov2001} with the two leading contenders often described as the disk-truncation and dynamic corona model \\citep[see][for a detailed study of these models]{done2002review,Beloborodov1999review}. To summarise, in the former, an increase in the reflection fraction is caused by the accretion disk penetrating deeper into a central hot corona thus receiving more illuminating hard photons. The presence of the disk in return offers more soft photons consequently cooling the plasma. For a purely thermal distribution of electrons, the greater the number of soft seed photons, the softer the power-law spectra.  The latter model by \\cite{beloborodov99} invokes bulk-motion of the corona above the accretion disk. If the corona is outflowing with mildly relativistic velocities, this would reduce the amount of hard photons hitting the disk, which in turn reduces $R$ and subsequently softens $\\Gamma$ as few reprocessed soft photons reach the outflow.  Recent evidence for the coronal plasma ejection model has come from a strong positive correlation between reflected X-ray flux and radio flux in the  black hole binary   Cygnus X-1 \\citep{miller2012cyg}.   We will show in \\S~4.6 that all evidence points toward the disk radius remaining stable during the HIS to SIS transition (see Figs.~11,12). This constancy effectively rules out the ``disk-truncation\" explanation for the $R-\\Gamma$ correlation. Instead, one may hypothesise whether the ``outflowing corona\" and light bending can be combined to explain the behavior so far detailed for \\j.  In the previous section we showed that the flux-flux behavior (Fig.~7) could be explained if early in the outburst (during the HIS) the corona was located relatively far ($\\sim10$\\rg)  from the black hole and thus  behaved according to  Regime 2 of \\citet[][see also \\S1.2]{Miniu04}. As the system evolved, the corona  collapsed to  a few~\\rg\\ and began to experience a higher level of light bending towards the disk (regime~1).  In this scenario, a potential gradient in the outflow velocity of the corona as a function of height could also explain  the behavior seen in Fig.~8. i.e. as the corona collapses from a large height (large outflow velocity; low $R$ and hard $\\Gamma$) the outflow velocity decreases ($R$ increases and $\\Gamma$ softens) until it becomes effectively static and the system transitions into the SIS.  We will expand upon this possible scenario in what follows and summarise our ideas in \\S~5.   \\begin{figure}[!t] \\hspace*{-0.5cm} \\centering {\\rotatebox{0}{{\\includegraphics[width=9.2cm]{fig_r_vs_gamma-eps-converted-to.pdf} }}} \\vspace{-0.5cm} \\caption{ The well know R-gamma relation appears present in the rising-LHS and during most of the HIS. However, just prior to the transition to the SIS and thereafter, this relation does not hold.  We discuss a potential explanation for this behavior in \\S~4.4. Only data with errors less than 50\\% their value are used in this figure. The black arrows show approximately the evolution of the system in time. }  \\vspace*{0.2cm} \\label{fig8} \\end{figure}   \\subsection{QPOs and Spectral States: A collapsing corona}    \\begin{figure*}[!t] \\hspace*{-0.35cm} \\centering {\\rotatebox{0}{{\\includegraphics[width=16cm]{fig_qpo_all_combined-eps-converted-to.pdf} }}} \\vspace*{-0.cm} \\caption{Top-left: HF QPO coherence as a function of reflection fraction.  The relation is well described by a linear function similar to that shown in the figure. Bottom-left: QPO frequency as a function of reflection fraction. The dashed line shows a relation, $f ({\\rm Hz}) = (102\\pm2)\\times {\\rm ln} [R \\times (3.07\\pm 0.05)]$, that fits these data. Right: QPO frequency as a function of disk surface ionisation. These figures show a clear link between the coherence/frequency of the QPO and the reflection fraction or level of disk ionisation,  which we subsequently  interpret as being linked to the size/position of the corona (see \\S~4.5). } \\label{fig9} \\end{figure*}    Throughout this work, we followed the selection made by Homan et al.~(2003) which roughly separates the outburst into six periods coinciding  with significant changes in both the hardness intensity diagram (Fig.~2) as well as in the shape of their power spectra\\footnote{In keeping with that work, the HIS is divided into two periods (P1 and P2). Here, we also add two extra periods which we have denoted as the LHS-rising and falling.} As highlighted in the introduction (\\S1.3), those authors demonstrate the presence of high-frequency (HF) variability in \\j, together with a  HFQPO which was shown to evolve in both frequency and coherence during the outburst.  The highest frequency which was reliably measured was at  $\\sim250$Hz in the SIS, with the frequency being much lower ($\\sim50$Hz) at the onset of the outburst.   In Fig.~9 (top-left), we show the presence of a strong (Pearson's $r=0.997$) positive relation between the reflection fraction and quality factor of the QPO ($Q-R$ relation). In order to create this figure, we have averaged the values of $R$ shown in Fig.~3  for each of the periods in question and used the values for the coherence provided by Homan et al.~(2003; Table~1). The bottom-left panel shows the frequency of the HF QPO as a function of $R$ ($f-R$ relation).  We also show in Fig.~9 (right), the QPO frequency as a function of disk (surface) ionisation parameter.     QPOs are notoriously difficult to explain and it is not our purpose to provide a  quantitative description of this phenomenon. However, it is worth stressing that most models \\citep[e.g.][]{Nowak1997, CuiZhangChen1998,  PsaltisBellonivanderKlis1999, StellaVietriMorsink1999, McKinneyqpo2012} strongly link the origin of QPOs with orbits and/or resonances in the inner accretion disk close to the black hole. Current models cannot fully explain, in a physical manner,  the range in coherence observed in various systems nor the manner in which the frequencies change with states.  To explain the  range in coherence observed in accreting neutron stars, \\citet{BarretOlive2006,BarretOlive2007} devised a toy model in which the  changes in $Q$ are driven by changes in the scale height of the disk. A small scale height gives rise to high coherence and vice versa. Expanding on this idea, it appears that at least for \\j, it is physical changes in the radius/size of the corona that give rise to changes in both quality factor $Q$ and QPO frequency. To illustrate this hypothesis, consider  Fig.~9 (bottom) together with our interpretation of the behavior displayed in Fig.~7 (\\S~4.3). Early in the outburst the frequency of the HF QPO  appeared at $\\sim55$Hz.  The Keplerian frequency at a given radius is $f ({\\rm Hz}) \\approx 3.2\\times 10^4  M^{-1} r^{-3/2}$, where $M$ and $r$ are in units of Solar mass and Gravitational radius respectively. Hence, during the brief P1 period, the HF QPO frequency is close to the orbital frequency at $\\sim28$\\rg\\ assuming a $4\\msun$ black hole and potentially moves to $\\sim15$\\rg\\ in the second half of the HIS. As the outburst continues and the corona continues to collapse, it is plausible that the frequency continues to increase (corresponding to $\\sim10$\\rg\\ in the SIS) eventually approaching a value that should be consistent with the  Keplerian frequency at the ISCO. The continued decrease in the size of the corona gives rise to the increase in coherence. In this scenario, the frequency of the QPO should relate to the size of the corona and thus would naturally increase as the corona collapses. The relationship between the QPO frequency and the surface ionisation parameter could be suggestive of an intrinsic relationship between the irradiation of the disk and its  magnetic field properties,  the latter which has recently been proposed as a possible means to  produce(low-frequency) QPOs \\citep[see e.g.][]{oneill2011qpo, Oishi2011qpo}. This possibility will be addressed in forthcoming work.     Finally, the relation between the coherence of the HF QPOs and the reflection fraction also leads to the interesting prediction  that HF QPOs should only be observed when $R\\gtrsim 0.4$ when the coherence is $Q>0$. This results is consistent with observations, with no HF QPO ever having been found in the LHS where $R\\lesssim1$.     \\subsection{Radio (jet) Emission and Reflection Fraction} \\label{jet}    \\citet{CorbelFender2004} presented a comprehensive analyses of the radio emission observed  during the outburst of \\j. In that work, the authors suggest that the transition between the HIS and SIS\\footnote{Referred to as the intermediate and steep powerlaw states respectively in \\citet{CorbelFender2004}.}  is associated with a massive radio  ejection event. The observations in the LHS was found to be consistent with the presence of a steady, compact jet, as is often seen in the LHS of black hole binaries \\citep[see e.g.][and references therein]{Fender2001jets,Fender091}. This potential ejection event during the HIS--SIS transition coincides with the time where we see a sharp jump in the reflection-fraction. In Fig.~10 we show the radio flux density\\footnote{Obtained directly from  Table~1 of \\citet{CorbelFender2004}.  We are using the flux densities at 4800~MHz  for all observations, except for the rising-LHS where this was not available. In this case,  we proceeded by averaging the values presented for 1384 and 2496~MHz.}  vs the reflection fraction calculated herein. At the time of the steady compact jet (in the LHS \\& HIS), the radio flux density increased by a factor of $\\sim5$ with no statistically significant change in $R$. However, immediately following the radio ejection in the SIS, the reflection fraction increases dramatically resulting in a bi-modality in the radio-flux density--reflection fraction plane. \\textit{This suggests an intimate link between the jet ejection site and the collapsing corona.}  Indeed, \\citet{beloborodov99}  predicts a link between radio jets and reflected flux, which was also   seen in Cygnus X-1 \\citep{miller2012cyg}.  At later stages in the outburst (HIS-P6, falling LHS), the measured radio flux density is likely a dominated by emission from the zone where the ejected plasma interacts with the ambient ISM surrounding the binary system.  Unfortunately, due to the low spatial resolution of the radio observations, this emission is not resolved from that due to any reformation of the steady jet close to the black hole.   \\begin{figure}[!t] \\hspace{-0.5cm} {\\rotatebox{0}{{\\includegraphics[width=9.cm]{fig_jet1-eps-converted-to.pdf} }}} \\vspace{-0.4cm} \\caption{ Radio flux density as a function of reflection fraction. The arrows show the direction of the outburst. There appears to be two distinct branches where the reflection fraction is either constant at $R\\sim0.5$ or at $R\\sim4$. These two branches corresponds to the LHS/HIS and SIS respectively. Radio data obtained from \\citet{CorbelFender2004}.   }\\vspace*{0.2cm} \\label{fig10} \\end{figure}      \\label{spin} \\subsection{Disk (inner) Radius and State Transition}   \\begin{figure*}[!t] \\vspace*{-0cm} \\centering {\\hspace{-0.0cm} \\rotatebox{0}{{\\includegraphics[ width=16.5cm]{fig_rin_hist3-eps-converted-to.pdf} }}} \\vspace*{-0.cm} \\caption{Left: The inner radius during the outburst are shown in blue. We only show measurements where the uncertainty in the radius  is $< 50\\%$ of its value. The dashed lines show the evolution of the PCA (grey) and HEXTE-A (green) count rate in a similar manner to Fig.~6 (bottom). Right: Distribution of inner radii having errors $\\leq50\\%$ its value. In red we show the standard normal distribution with a mean radius of 1.65\\rg\\ and standard deviation of  0.08\\rg. } \\label{fig11} \\end{figure*}    \\begin{figure*}[] \\vspace*{-0.4cm}  \\centering {\\hspace{-0.0cm} \\rotatebox{0}{{\\includegraphics[ width=8.cm]{fig_mcmc_link1-eps-converted-to.pdf} }}} {\\hspace{-0cm} \\rotatebox{0}{{\\includegraphics[ width=8.cm]{fig_mcmc_chain2-eps-converted-to.pdf} }}}  \\vspace*{-0.cm} \\caption{MCMC results for  the inner radius (a proxy for spin) obtained from the simultaneous fit to the first ten spectra in the HIS. Left: Figure tracing 50 out of a total of 170 ``walkers\" during their random walk. The figure shows that the various chains converge pretty quickly indicating efficient sampling. The inset shows a close up of the first 20,000. steps.  Right: Full MCMC containing all 170 walkers, after having the first 5,000 elements ``burnt-in\". For clarity, we only show every 1000th element of the chain.}  \\vspace*{0.2cm}  \\label{fig12} \\end{figure*}    As has been discussed throughout this paper, a popular explanation for state transitions is a radial variation in the extent of  the  accretion disk. This model has been highly successful in part due to its flexibility and ease in which it can explain the ``weak\" reflection fraction ($R<1$) often found in the LHS. However, over the past few years it has consistently been shown that in the luminous phases of the LHS -- at least above $\\sim1\\times 10^{-3}~L_{\\rm Edd}$\\footnote{Contrast this with the broadband analyses of GX~339-4 presented by \\citet{tomsick09gx} where the authors find clear evidence for the recession of the accretion disk beginning  only at Eddington luminosities below $\\sim\\times10^{-3}~L_{\\rm Edd}$.  } -- the disk does not appear to be truncated away from the radius of the ISCO \\citep{Millergxlhs2006,  MillerHomanMiniutti2006, miller08gx, reisj1118, reislhs, Reynold2011swift, Waltonreis2012}.   Figure~11 (left) shows that the present work can statistically rule out a disk being truncated further than $\\sim3$\\rg\\ even in the LHS. During the brighter, intermediate states, we constrain this radius to $\\sim1.65$\\rg. This adds support to the idea that the inner disc radius remains roughly constant throughout the  LHS--HIS--SIS state transitions in black hole binaries.    Where we have not been able to constrain the radius, this has largely been due to the data quality (the falling phase of the LHS is inherently less luminous)  as well as the fact that reflection is intrinsically weaker in the LHS. The strongest reflection features are expected in the intermediate states where the disk receives a larger fraction of the hard X-ray emission \\citep[see e.g.][]{hiemstra1652}. Note also that despite the comparatively  low spectral resolution afforded by the \\rxte -PCA (18\\% FWHM energy resolution at 5.9\\kev)~\\citet{WilmsNowak2006cyg} showed that this instrument  can indeed resolve line-widths down to a least  $\\sigma \\sim0.3$\\kev. Higher resolution observations of \\j\\ early in the outburst showed that, when modelled with a \\gaussian, the Fe\\ka\\ emission line is consistent with having a width of~$\\sigma \\sim~1.1$\\kev~\\citep[Table~4 in][]{Waltonreis2012}\\footnote{The original analyses of this \\xmm\\ dataset performed by \\citet{miller02j1650} included an extra smeared edge component at $\\sim6.8$\\kev\\ which resulted in the \\gaussian\\ having a width of only $\\sim250$\\ev. }.  Assuming  the stable radius shown in Fig.~11 for the two intermediate states is indeed the radius of the ISCO, we find, using the relationship between ISCO and black hole spin of \\citet{Bardeenetal1972}, a dimensionless spin parameter of  $0.977^{+0.006}_{-0.007}$ consistent with the value found in  detailed analyses of single, high quality data obtained with \\xmm\\ ($0.84\\leq a \\leq 0.98$; \\citealt{Waltonreis2012}) or  \\bepposax\\ ($a \\gtrsim 0.93$; \\citealt{MiniuttiFabianMiller2004j1650}\\footnote{Spin converted from the lower limit on the inner radius of $\\sim2.1$\\rg.}).   As a test of the robustness of this result, we have performed a joint fit to the first ten observations in the HIS. We used the same base model as before with each individual observation having their own set of parameters -- disk emissivity index, temperature, normalisation and  parameter,   power-law index, as well as the normalisation of the powerlaw and reflection component. However, this time we  forced the inner radius in the various observations to be a global parameter thus assuming a constant value.   This simultaneous fit contains a total of 81 free parameters\\footnote{The sheer number of free parameters and computational time required to do  $\\chisq$ fitting as well as the MCMC analyses described in what follows drove the need to constrain this analyses to only 10 observations as opposed to all 116.} and with this comes a high chance of mistaking a local minima in $\\chisq$ space for the global best fit. In order to address these limitations, we proceeded by minimised the fit using standard $\\chisq$ fitting techniques within \\xspec\\ until a reasonable fit was produced ($\\chisq/\\nu <2$) at which point we halted the minimisation\\footnote{The actual quality of the fit at this time was $\\chisq/\\nu = 2059.7/1149$.} and proceeded with Monte Carlo Markov Chain (MCMC) analysis.  We employed the MCMC procedure described in \\citet{mcmc2012} (code  found at \\href{http://danfm.ca/emcee/}{http://danfm.ca/emcee/}) and implemented in the \\xspec  spectral fitting package by Jeremy Sanders (\\xspec implementation described in \\href{https://github.com/jeremysanders/xspec_emcee}{https://github.com/jeremysanders/xspec\\_emcee}). MCMC techniques have been successfully used to address similar problems in constraining the black hole spin of NGC~3783 \\citep{Reynolds20123783} as well as in modelling the kinematics of the microquasor XTE~J1550--564 \\citep{SteinerMcClintock2012jet1550}.     We added  a 5\\% random perturbation to all the parameters in the fit  described above, and increased the value of the  inner radius from the  starting  value of $\\sim1.6$\\rg\\ to 2.5\\rg\\ in order to guarantee that the chain could freely converge to the global minimum.  We used a total of 170 ``walkers\", each iterated (``walking\")  10,000 times. Figure~12 (left) shows the evolution of the walk in inner radius for 50 randomly selected walkers. It is clear that the walkers converge to the same value  efficiently. Nonetheless, in order to be conservative we have ignored (``burned-in\") the first 5000 elements of each chain and show on the right the full MCMC chain for the radius which is clearly well behaved with a peak distribution at $1.66\\pm0.01$\\rg\\ (s.d.), in excellent  agreement with the results in Fig.~11."
    },
    "1208/1208.1150.txt": {
        "abstract": "We construct full broadband models in an analysis of {\\sl Suzaku} observations of nearby Seyfert 1 AGN ($z\\leq0.2$) with exposures $>50$\\,ks and with greater than 30000 counts in order to study their iron line profiles. This results in a sample of 46 objects and 84 observations. After a full modelling of the broadband {\\sl Suzaku} and {\\sl Swift}-BAT data (0.6-100\\,keV) %, curvature which may be introduced as a consequence of warm absorbers has a significant affect upon the spectrum at energies $>3$\\,keV and notably in the FeK region. Accounting for the warm absorber and using the broadband data allows us to assess the true extent to which relativistic emission is present and as a result allows a more robust measurement of accretion disc and black hole parameters. we find complex warm absorption is present in $59\\%$ of the objects in this sample which has a significant bearing upon the derived Fe\\,K region parameters. Meanwhile $35\\%$ of the 46 objects require some degree of high column density partial coverer in order to fully model the hard X-ray spectrum. We also find that a large number of the objects in the sample ($22\\%$) require high velocity, high ionization outflows in the Fe\\,K region resulting from Fe\\,{\\rm XXV} and Fe\\,{\\rm XXVI}. A further four AGN feature highly ionized Fe\\,K absorbers consistent with zero outflow velocity, making a total of 14/46 ($30\\%$) AGN in this sample showing evidence for statistically significant absorption in the Fe\\,K region.  Narrow Fe\\,K$\\alpha$ emission from distant material at 6.4\\,keV is found to be almost ubiquitous in these AGN. Examining the 6-7\\,keV Fe\\,K region we note that narrow emission lines originating from Fe\\,{\\rm XXV} at 6.63-6.70\\,keV and from Fe\\,{\\rm XXVI} at 6.97\\,keV are present in $52\\%$ and $39\\%$ of objects respectively.  Our results suggest statistically significant relativistic Fe\\,K$\\alpha$ emission is detected in 23 of 46 objects ($50\\%$) at $>99.5\\%$ confidence, measuring an average emissivity index of $q=2.4\\pm0.1$ and equivalent width $EW=96\\pm10$\\,eV using the \\textsc{relline} model. When parameterised with a Gaussian we find an average line energy of $6.32\\pm0.04$\\,keV, $\\sigma_{\\rm width}=0.470\\pm0.05\u00a4$\\,keV and $EW=97\\pm19$\\,eV. Where we can place constraints upon the black hole spin parameter $a$, we do not require a maximally spinning black hole in all cases. ",
        "introduction": "The analysis of the X-ray spectra of AGN can reveal information regarding the inner regions of the accretion disc, the AGN environment as a whole and subsequently the super massive black hole (SMBH) at its heart. It was suggested by Fabian et al. (1989) that emission occurring from the very inner regions of the accretion disc may be visible and subsequently broadened by Doppler motions and relativistic affects. The majority of AGN spectra show narrow line emission from neutral iron at 6.4\\,keV (Fe\\,K$\\alpha$) likely originating from distant material e.g. the torus or the outer regions of the accretion disc (Krolik \\& Kallman 1987; Nandra 2006), particularly strong due to the high abundance and fluorescent yield of iron. If Fe\\,K$\\alpha$ emission additionally arises from material close to the central SMBH it will become relativistically broadened (Fabian et al. 1989; Laor 1991), producing a both blue and red-wings to the traditionally narrow line profile.  In some AGN spectra this relativistic Fe\\,K$\\alpha$ emission may be strong enough to be observed allowing its shape and strength measured using disk-line emission models such as \\textsc{laor}, \\textsc{kyrline}, \\textsc{kerrdisk} and \\textsc{relline} (Laor 1991; Dov\\v{c}iak et al. 2004; Brenneman \\& Reynolds 2006; Dauser et al. 2010). The application of these models allows properties such as the inclination and emissivity index of the disc to be measured in addition to the typical inner radius of emission and in some cases the spin of the central SMBH (Nandra et al. 2007; Patrick et al. 2011a). Gaining information regarding the distribution of SMBH spins is an essential tool in aiding our understanding of galaxy evolution and distinguishing between models such as prolonged or chaotic accretion and also the effect of mergers upon the SMBH spin (Hughes \\& Blandford 2003; Volonteri et al. 2005; King \\& Pringle 2007; Rezzolla et al. 2008). A spin distribution skewed towards higher values ($a\\sim0.998$) would suggest prolonged accretion, whereas low SMBH spin ($a\\sim0$) would indicate chaotic accretion models are more appropriate. In addition to this, the magnetic extraction of BH rotational energy through the Blandford-Znajek effect (Blandford \\& Znajek 1977) could cause a reduction in the spin (i.e. towards zero) of the SMBH in some AGN (Berti \\& Volonteri 2008).  %In addition to these scenarios, it has been suggested that a stable configuration can be obtained with a {\\sl retrograde} spin whereby the direction in which the BH spins is opposite to that of the surrounding accretion disc (King et al. 2005). There have been tentative indications of retrograde SMBHs, to date only NGC 3783 has had spin very tentatively measured at $a<-0.04$ (Patrick et al. 2011b) using the \\textsc{relline} line emission model of Dauser et al. (2010), although it should be stressed that this could also be explained by the truncation of the accretion disc at $\\gtrsim6\\,R_{\\rm g}$ i.e. short of the innermost stable circular orbit for a Schwarzschild BH. If a retrograde spin is also an efficient scenario for the production of radio jets in AGN (Garofalo 2009), then it is expected that some radio-loud AGN would show indications of an accretion disc truncated further out than the innermost stable circular orbit (ISCO) at typically $R_{\\rm in}\\sim9\\,R_{\\rm g}$ (i.e. $R_{\\rm in}$ for a maximal retrograde spin $a=-0.998$).  Recent publications have made steps towards making spin estimates of SMBHs in a variety of AGN, including those which feature low levels of intrinsic absorption thereby offering the simplest spectrum to analyse, avoiding complications regarding the degree of spectral curvature introduced with warm absorbing zones (Miniutti et al. 2009; Schmoll et al. 2009; Patrick et al. 2011a; Emmanoulopoulos et al. 2011). More complex AGN spectra have also been analysed and revealed further spin estimates e.g. MCG--06-30-15 (Miniutti et al. 2007; Patrick et al. 2011b), Mrk 79 (Gallo et al. 2011) and NGC 3783 (Brenneman et al. 2011; Patrick et al. 2011b). However, as discussed in Patrick et al. (2011a), the estimated SMBH spin is highly model dependent and strongly related to the treatment of features such as the soft excess or any intrinsic absorbing zones. Assuming a Comptonization origin of the soft excess results in a range of low to intermediate spins, whereas using a high degree of relativistic blurring to smooth the discrete soft emission lines into a continuum typically forces the spin to near maximal values and requires very high disc emissivities.  This paper includes a sample of AGN from the public {\\sl Suzaku} archive of all observations of Seyfert 1 AGN with total exposures $>50$\\,ks and more than 30\\,000 counts in order to increase the likelihood of detection and broadened emission from the inner regions if it is present. {\\sl Suzaku} is the ideal instrument with which to do this work since it allows us to gather both soft and hard X-ray data simultaneously using the X-ray Imaging Spectrometer (Koyama et al. 2007) and Hard X-ray Detector (Takahashi et al. 2007) detectors which, when combined with further non-simultaneous hard X-ray data from {\\sl Swift} BAT, gives a broad energy bandpass of 0.6-100.0\\,keV. The crucial difference being that we can obtain data regarding the strength of the Compton reflection hump at $\\sim30$\\,keV (George \\& Fabian 1991) which is beyond the capabilities of other current X-ray observatories. Only with hard X-ray data can the strength of the reflection component be appropriately constrained and hence its contribution to the Fe\\,K region assessed prior to attempting to determine broadening in the Fe\\,K region and eventually estimates upon SMBH spin.  This is the final paper in a series of three in which a methodical and relatively uniform approach has been taken in an attempt to constrain accretion disc properties and SMBH spin from the Fe\\,K regions from an analysis of the X-ray spectra of Seyfert 1 AGN. In Patrick et al. (2011a) a small sample of six 'bare' Seyfert 1 AGN (i.e. those featuring low intrinsic absorption: Ark 120, Fairall 9, MCG--02-14-009, Mrk 335, NGC 7469 and SWIFT J2127.4+5654) was analysed, finding that narrow ionized emission lines such as Fe\\,{\\rm XXVI} are relatively common (4/6 objects) while the emissivity index of the accretion disc indicated that no strongly centrally concentrated emission was required to model any relativistic broadening in the Fe\\,K region with an average of $q\\sim2.3\\pm0.2$.  Patrick et al. (2011b) undertook and analysis of high quality, long exposure ($>200$\\,ks) observations of Seyfert 1 AGN with {\\sl Suzaku} (Fairall 9, MCG--06-30-15, NGC 3516, NGC 3783 and NGC 4051) making use of the full 0.6-100.0\\,keV bandpass in order to fully account for any warm absorber component to give the best possible opportunity to make estimates upon SMBH spin, also finding a low to moderate average emissivity index of $q\\sim2.8\\pm0.2$.  %Spin estimates to date within Patrick et al. (2011a, 2011b) are: Ark 120 ($a<0.94$), Fairall 9 ($a=0.67^{+0.10}_{-0.11}$), MCG--02-14-009 ($a<0.88$), MGC--6-30-15 ($a=0.49^{+0.20}_{-0.12}$), Mrk 335 ($a=0.70^{+0.12}_{-0.01}$), NGC 3783 ($a<0.32$), NGC 4051 ($a<0.94$), NGC 7469 ($a=0.69^{+0.09}_{-0.09}$) and SWIFT J2127.4+5654 ($a=0.70^{+0.10}_{-0.14}$). Other authors have also made some progress towards making spin estimates on Seyfert 1 AGN e.g. Miniutti et al. (2009), Schmoll et al. (2009), Nardini et al. (2011), Gallo et al. (2011), Emmanoulopoulos et al. (2011) and Brenneman et al. (2011).  The main aim of this paper is to assess the properties and total percentage of AGN which have been observed with {\\sl Suzaku} that show evidence for broadened line emission in the Fe\\,K region from the inner regions of the accretion disc resulting from an analysis of the broad-band X-ray spectrum. In this paper we expand our broad-band spectral analysis with {\\sl Suzaku} to all the currently archived type 1-1.9 AGN, which have at least a 50\\,ks total exposure and 30000 XIS band counts, sufficient for a broad band spectral analysis. This enables us to measure the iron line and reflection properties of a more substantial sample of 46 type 1 AGN, allowing the overall properties of the iron line and accretion disc to be investigated. We also aim to investigate {\\sl Suzaku's} view of ionised emission and absorption lines in the Fe\\,K region and the occurrence or warm absorbers, highly ionised absorbers and partially covering absorbers and the subsequent effects upon the AGN X-ray spectrum.  %This paper makes use of previous results, models and methods in order to produce as uniform an analysis as possible from which the typical properties of the accretion disc and Fe\\,K region can be estimated.   \\begin{table*} \\caption{The {\\sl Suzaku} Seyfert sample} \\begin{tabular}{l c c c c} \\hline Object  & RA (J2000) & Dec (J2000) & Redshift & $N_{H}$ (Gal) ($10^{22}$cm$^{-2}$) \\\\ \\hline %1H 0323+342 & 03 24 41.2 & +34 10 45.9 & 0.0610 & 0.1270 \\\\ 1H 0419--577 & 04 26 00.8 & --57 12 00.4 & 0.1040 & 0.0126 \\\\ %1H 0707-495 & 07 08 41.5 & --49 33 06.4 & 0.0406 & \\\\ 3C 111 & 04 18 21.3 & +38 01 35.8 & 0.0485 & 0.2910 \\\\ 3C 120 & 04 33 11.1 & +05 21 15.6 & 0.0330 & 0.1060 \\\\ %3C 326 & 15 52 09.1 & +20 05 35.8 & 0.0895 & \\\\ 3C 382 & 18 35 03.4 & +32 41 46.8 & 0.0579 & 0.0714 \\\\ 3C 390.3 & 18 42 09.0 & +79 46 17.1 & 0.0561 & 0.0347 \\\\ 3C 445 & 22 23 49.5 & --02 06 12.9 & 0.0559 & 0.0559 \\\\ %4C 73.08 & 09 49 45.9 & +73 14 23.1 & 0.0581 & \\\\ 4C 74.26 & 20 42 37.3 & +75 08 02.4 & 0.1040 & 0.1160 \\\\ Ark 120 & 05 16 11.4 & --00 08 59.4 & 0.0327 & 0.0978 \\\\ Ark 564 & 22 42 39.3 & +29 43 31.3 & 0.0247 & 0.0534 \\\\ Fairall 9 & 01 23 45.8 & --58 48 20.5 & 0.0470 & 0.0316 \\\\ IC 4329A & 13 49 19.3 & --30 18 34.0 & 0.0161 & 0.0461 \\\\ IRAS 13224--3809 & 13 25 19.4 & --38 24 52.7 & 0.0658 & 0.0534 \\\\ MCG--02-14-009 & 05 16 21.2 & --10 33 41.4 & 0.0285 & 0.0924 \\\\ MCG--02-58-22 & 23 04 43.65 & --08 41 08.6 & 0.0649 & 0.0291 \\\\ MCG--05-23-16 & 09 47 40.2 & --30 56 55.9 & 0.0085 & 0.0870 \\\\ MCG--06-30-15 & 13 35 53.8 & --34 17 44.1 & 0.0077 & 0.0392 \\\\ MCG+8-11-11 & 05 54 53.6 & +46 26 21.6 & 0.0205 & 0.1840 \\\\ MR 2251--178 & 22 54 05.8 & --17 34 55.0 & 0.0640 & 0.0640 \\\\ Mrk 79 & 07 42 32.8 & +49 48 34.8 & 0.0222 & 0.0527 \\\\ Mrk 110 & 09 25 12.9 & +52 17 10.5 & 0.0353 & 0.0130 \\\\ Mrk 205 & 12 21 44.0 & +75 18 38.5 & 0.0708 & 0.0280 \\\\ Mrk 279 & 13 53 03.5 & +69 18 29.6 & 0.0305 & 0.0152 \\\\ Mrk 335 & 00 06 19.5 & +20 12 10.5 & 0.0258 & 0.0356 \\\\ Mrk 359 & 01 27 32.5 & +19 10 43.8 & 0.0174 & 0.0426 \\\\ Mrk 509 & 20 44 09.7 & --10 43 24.5 & 0.0344 & 0.0344 \\\\ Mrk 766 & 12 18 26.5 & +29 48 46.3 & 0.0129 & 0.0178 \\\\ Mrk 841 & 15 04 01.2 & +10 26 16.2 & 0.0364 & 0.0222 \\\\ %Mrk 1239 & 09 52 19.1 & --01 36 43.5 & 0.0199 & 0.0369 \\\\ %NGC 1194 & 03 03 49.1 & --01 06 13.5 & 0.0136 & \\\\ NGC 1365 & 03 33 36.4 & --36 08 25.5 & 0.0055 & 0.0134 \\\\ NGC 2992 & 09 45 42.1 & --14 19 35.0 & 0.0077 & 0.0487 \\\\ NGC 3147 & 10 16 53.7 & +73 24 02.7 & 0.0093 & 0.0285 \\\\ NGC 3227 & 10 23 30.6 & +19 51 54.2 & 0.0039 & 0.0199 \\\\ NGC 3516 & 11 06 47.5 & +72 34 06.9 & 0.0088 & 0.0345 \\\\ NGC 3783 & 11 39 01.7 & -37 44 18.9 & 0.0097 & 0.0991 \\\\ NGC 4051 & 12 03 09.6 & +44 31 52.8 & 0.0023 & 0.0115 \\\\ %NGC 4138 & 12 09 29.8 & +43 41 07.1 & 0.0030 & \\\\ NGC 4151 & 12 10 32.6 & +39 24 20.6 & 0.0033 & 0.0230 \\\\ %NGC 4395 & 12 25 48.9 & +33 32 48.9 & 0.0011 & \\\\ NGC 4593 & 12 39 39.4 & --05 20 39.3 & 0.0090 & 0.0189 \\\\ NGC 5506 & 14 13 14.9 & --03 12 27.3 & 0.0062 & 0.0408 \\\\ NGC 5548 & 14 17 59.5 & +25 08 12.4 & 0.0172 & 0.0155 \\\\ NGC 7213 & 22 09 16.3 & --47 09 59.8 & 0.0058 & 0.0106 \\\\ NGC 7314 & 22 35 46.2 & --26 03 01.7 & 0.0048 & 0.0150 \\\\ NGC 7469 & 23 03 15.6 & +08 52 26.4 & 0.0163 & 0.0445 \\\\ PDS 456 & 17 28 19.8 & --14 15 55.9 & 0.1840 & 0.1960 \\\\ PG 1211+143 & 12 14 17.7 & +14 03 12.6 & 0.0809 & 0.0274 \\\\ RBS 1124 & 12 31 36.4 & +70 44 14.1 & 0.2080 & 0.0152 \\\\ %SWIFT J0318.7+6828 & 03 18 19.0 & +68 29 31.4 & 0.0901 & \\\\ %SWIFT J0444.1+2813 & 04 44 09.0 & +28 13 00.5 & 0.0113 & \\\\ SWIFT J2127.4+5654 & 21 27 45.0 & +56 56 39.7 & 0.0144 & 0.7650 \\\\ TON S180 & 00 57 19.9 & --22 22 59.1 & 0.0620 & 0.0136 \\\\ \\hline \\end{tabular} \\label{tab:sample} \\end{table*} ",
        "conclusions": "Based upon an analysis of all publically available {\\sl Suzaku} observations of Seyfert 1 AGN with observations longer than 50\\,ks and greater than 30000 XIS counts, we conclude the following: \\begin{enumerate} \\item The majority ($59\\%$) of AGN in this sample feature complex warm absorption which has a significant affect upon the Fe\\,K region and any accretion disc parameters derived from it. The use of the full 0.6-100.0\\,keV  broad-band data is therefore essential prior to any attempt to use relativistic line profile models as a diagnostic for the inner regions of the accretion disc. The mean photon index of the 46 objects on the sample is $\\Gamma=1.82\\pm0.03$. %Failure to do so can result in improper modelling of the continuum and insufficient spectral curvature resulting in excessively broad residuals. Fitting relativistic line emission models or blurred reflection convolutions to these residuals then forces the emissivity and spin estimates to very high near maximal values. \\item Absorption in the Fe\\,K region due to highly ionized gas producing absorption features from Fe\\,{\\rm XXV} and Fe\\,{\\rm XXVI} are relatively common in AGN (30\\%), most of which (71\\%) are outflowing at high velocities. While a large fraction of the detected highly ionised winds are outflowing, there may be a larger number of low velocity winds which are not detected due to the possible presence of ionised Fe\\,{\\rm XXV} and Fe\\,{\\rm XXVI} emission lines. The additional curvature added to the region through modelling with an appropriate \\textsc{xstar} grid, while subtle, has a notable effect upon the strength of any broad residuals which may remain below 6.4\\,keV and be interpreted as strong relativistic emission. \\item A partial covering geometry is required in 35\\% of all objects in the sample. These high column density zones primarily affect the hard X-ray spectrum above 7\\,keV although reducing the strength of broad residuals in the Fe\\,K region rather than removing them entirely. %Measured accretion disc parameters and SMBH spin estimates remain relatively unchanged however. \\item Narrow ionized emission in the Fe\\,K region from Fe\\,{\\rm XXV} and Fe\\,{\\rm XXVI} are relatively common in these AGN, featuring in 24/46 and 18/46 of objects respectively. Of these AGN, 10/46 feature both Fe\\,{\\rm XXV} and Fe\\,{\\rm XXVI} emission. These lines are found to be much more common compared to an {\\sl XMM-Newton} survey by Nandra et al. (2007) despite possible interplay with the blue-wing of more sophisticated relativistic line emission models which could reduce the number of narrow line detections. \\item Examining the Fe\\,K region after a complete modelling of the broad-band 0.6-100.0\\,keV spectrum and all required absorption zones yields a range of weak to moderate strength broad residuals below 6.4\\,keV. We find that 26/46 ($56\\%$) of the objects in this sample require some degree of relativistic line emission in the Fe\\,K region at $90\\%$ confidence and 23/46 (50\\%) at $>99.5\\%$ confidence. \\item These broad residuals are well fit with the \\textsc{relline} model and yield an average broad line strength of $EW=96\\pm10$\\,eV for the total of 26 objects. The line energy and $\\sigma_{\\rm width}$ of the broad residuals when modelled with a Gaussian are consistent with Nandra et al. (2007). \\item We estimate an average emissivity index of the accretion disc of $q=2.4\\pm0.1$, suggesting that emission from the accretion disc responsible for relativistic lines is not extremely centrally concentrated when purely the line profile in the Fe\\,K region is used as a diagnostic. The majority of the line flux therefore occurs from the blue-wing of the line profile with emission being insufficiently close to $R_{\\rm ISCO}$ as to redshift a significant proportion of the X-ray flux into a strong red-wing. We also measure an average disc inclination of $i=33^{\\circ}\\pm2^{\\circ}$ and inner radius of emission $R_{\\rm in}=(21\\pm6)\\,R_{\\rm g}$. \\item With the assumption that the inner radius of the accretion disc ($R_{\\rm in}$) extends down to the innermost stable circular orbit ($R_{\\rm ISCO}$), loose constraints upon the SMBH spin parameter $a$ can be made. The relativistic line emission profiles are sufficiently distinguished in 11/46 objects to place upper or lower bounds on the spin. After a broad-band analysis we make the following estimates: Ark 120, $a<0.94$; Fairall 9, $a=0.60^{+0.19}_{-0.63}$; MCG--02-14-009, $a<0.88$; MCG--05-23-16, $a<0.50$; MCG--06-30-15, $a=0.49^{+0.20}_{-0.12}$; Mrk 79, $a<0.80$; Mrk 335, $a=0.70^{+0.12}_{-0.01}$; NGC 3516, $a<0.30$; NGC 3783, $a<0.24$; NGC 7469, $a=0.69^{+0.09}_{-0.09}$ and SWIFT J2127.4+5654 with $a=0.70^{+0.10}_{-0.14}$. Under the assumption that $R_{\\rm in}=R_{\\rm ISCO}$, a maximally rotating SMBH is ruled out in each of these 11 objects. \\end{enumerate}"
    },
    "1208/1208.6102_arXiv.txt": {
        "abstract": "We use Renewal Theory for the estimation and interpretation of the flare rate from the \\textit{Geostationary Operational Environmental Satellite} (GOES) soft \\Xr flare catalogue. It is found, that in addition to the flare rate variability with the solar cycles, a much faster variation occurs. The fast variation on time scales of days and hours down to minute scale appears to be comparable with time intervals between two successive flares (waiting times). The detected fast non-stationarity of the flaring rate is discussed in the framework of the previously published stochastic models of the waiting time dynamics. ",
        "introduction": "The phenomena related to energy transformation and release on the Sun, possibly are of highest importance for modern solar astrophysics. Solar flaring, as a steady process of energy release, plays a central role and has been drawing the attention of the scientific community for decades.  In this work, we focus on the statistical properties of so-called solar flares \\textit{waiting time(s)} (wt), \\textit{i.e.} the interval between two flares close in time. The data are provided by the \\textit{Geostationary Observational Environmental Satellites} (GOES) in the soft \\Xr band. This catalogue is chosen as the longest available record of uninterrupted observations. The \\Ha flare  records from the NGDC-NOAA catalogue$\\footnote{http://www.ngdc.noaa.gov/stp/solar/solarflares.html}$ are shorter, and used to emphasise the invariance of the reported results.  Flare waiting times statistics has been extensively debated in the literature. \\inlinecite{Boffetta} fitted the waiting time probability density function (pdf)\\footnote{We use the mathematically strict definition of the \\textit{density} as the differential quantity with respect to the \\textit{distribution}, which is often used in a misleading way instead.} with a power-law in the range between $6$ and $67$ hours. This estimate was implemented using a single record over $20$ years. \\inlinecite{Wheat-II} considered the variation of the mean flaring rate with the solar cycle, and proposed the model of the time-dependent Poisson process (see also \\opencite{Moon2001}). In turn, this model gives a power-law-like \\pdf only in the limit \\cite{Wheat-final}. However, \\inlinecite{Lepreti} argued for a local departure of the \\wt series from the Poisson process, \\textit{i.e.} \"memory\" had been detected in the data.  We apply an alternative statistical description to the methods that had been used in the articles cited above. This method is equivalently unambiguous in describing the stochastic processes. The flare rate is estimated explicitly from the waiting time pdf. Such an approach has the significant advantage of linking the waiting time of the flare and its rate, which is the instantaneous probability of the flare per unit of time.  The paper is organised as follows. In Section~\\ref{math} the rather simple mathematics used in the paper is summarised, considered for the paper's consistency to avoid referring the reader to specific literature. The main results are presented Section~\\ref{data_analysis}, and interpreted in Section~\\ref{discussion}. The conclusion and final remarks are presented in Section~\\ref{conclusion}. ",
        "conclusions": "The fast variation on time scales from minutes to hours of the flare rate is a highlight among other findings in this work. The large-scale variation with the solar cycle exhibits a complexity that can be coped with by a phase-wise splitting of the data. However, the fast non-stationarity during the waiting time appears to be an intrinsic feature of the flaring dynamics, which requires a further elaboration of the present solar flaring models.  It is worth mentioning that we see somewhat intermediate character of the reported results with respect to the works by \\inlinecite{Wheat-final}, \\inlinecite{Boffetta} and by \\inlinecite{Lepreti}. The \"memory\" revealed in the data rejects models with Poisson statistics. From another perspective, a simple power-law fit very likely fails for the records relevant to a specific solar cycle phase.  The piecewise-constant Poisson model had inspired us to consider solar cycle phases separately. We support the argument of this model that solar cycle phases having quite inhomogeneous statistical properties that are to be considered jointly. In other words, the separate consideration of the solar cycles should be a cornerstone of a realistic modelling of the solar flaring activity.  On the other hand, our findings exclude the presence of the timescale of true rate constancy (with GOES catalog precision, at least). Even if it exists, it would be very small, and may be hard to detect.  The application of Renewal Theory emphasises the importance of the short (energetically small) events in the pdf, whose dynamics had not been pointed out in the literature: partially due to removal from the dataset of the $B$-class flares, and partially due to systematic fitting solely the tails of the \\pdf's. This brings us to problems similar to those that had arisen in the context of coronal heating by nanoflares, where, perhaps, the most significant effects are at the limit of the instrumental noise.  \\begin{acks} We thank the GOES teams at NOAA and SIDC for data management and availability and Christoph Keller for useful comments and discussions. M.M. acknowledges the support of the Italian Space Agency (ASI) and COST Action ES0803. Mrs. S. Fabrizio (INAF-OATS) is gratefully acknowledged for careful proofreading. \\end{acks}"
    },
    "1208/1208.6334_arXiv.txt": {
        "abstract": "In order to find more examples of the elusive high-redshift molecular absorbers, we have embarked on a systematic discovery program for highly obscured, radio-loud ``invisible AGN'' using the VLA Faint Images of the Radio Sky at Twenty centimeters (FIRST) radio survey in conjunction with Sloan Digital Sky Survey (SDSS) to identify 82 strong ($\\geq$ 300 mJy) radio sources positionally coincident with late-type, presumably gas-rich galaxies. In this first paper, the basic properties of this sample are described including the selection process and the analysis of the spectral-energy-distributions (SEDs) derived from the optical (SDSS) $+$ near-IR (NIR) photometry obtained by us at the Apache Point Observatory 3.5m. The NIR images confirm the late-type galaxy morphologies found by SDSS for these sources in all but a few (6 of 70) cases (12 previously well-studied or misclassified sources were culled). Among 70 sources in the final sample, 33 show galaxy type SEDs, 17 have galaxy components to their SEDs, and 20 have quasar power-law continua. At least 9 sources with galaxy SEDs have $K$-band flux densities too faint to be giant ellipticals if placed at their photometric redshifts. Photometric redshifts for this sample are analyzed and found to be too inaccurate for an efficient radio-frequency absorption line search; spectroscopic redshifts are required. A few new spectroscopic redshifts for these sources are presented here but more will be needed to make significant progress in this field. Subsequent papers will describe the radio continuum properties of the sample and the search for redshifted \\ion{H}{1} 21~cm absorption. ",
        "introduction": "While most strong radio sources are identified with luminous elliptical galaxies or optical/UV point sources (quasars), a small number of these sources lie behind obscuring screens of gas and dust becoming invisible to us at optical/UV and possibly also at near-IR wavelengths.  Identifying highly obscured, ``invisible AGN'' can lead to the detection and study of highly-redshifted radio-frequency atomic and molecular absorption line systems, which can provide important information about the early universe unobtainable in other ways. Radio absorption lines provide unparalleled velocity resolution and detailed physical diagnostics. And, where the background radio source is spatially-resolved, very-long baseline interferometry maps in the \\ion{H}{1} absorption line can provide the best spatial resolution ($\\sim$ 10 pc at $z$=1--2) of the gas distribution in distant galaxies; magnetic field structures also can be mapped using Faraday rotation measures of a polarized background source \\citep{gae07}, or using Zeeman splitting in the 21 cm line itself in the future. Atomic plus molecular absorption can provide much additional information. First, molecules trace the gas that supplies star formation so that the strengths and ratios of their absorption lines can be used to determine the physical conditions (e.g., temperature, pressure, density, chemical abundance and magnetic field strength) of the locations where most stars have formed over cosmic time. Second, the combination of various \\ion{H}{1} and molecular absorption/emission lines makes it possible to probe whether and by how much dimensionless fundamental physical constants have changed as the universe evolves.  So far there are about 80 \\ion{H}{1} 21~cm absorbers detected at $z>0.1$. Only five of them are molecular absorbers\\footnotemark. The first system, PKS~1413+135, was detected in \\ion{H}{1} absorption in 1992 at $z = 0.25$ (Carilli, Perlman \\& Stocke 1992), in CO absorption in 1994 (Wiklind \\& Combes) and subsequently in nearly a dozen species (Wiklind \\& Combes 1997). This source is a BL Lac object residing in an edge-on spiral galaxy and the nucleus is highly obscured (A$_{\\textrm {\\scriptsize V}} >$ 20; Stocke et al.1992; Carllli et al. 1992). Then B3~1504+377 was found at $z=0.67$ (Wiklind \\& Combes 1996a), which is also an intrinsic molecular absorption system. The other three molecular absorbers are all in gravitational lens systems: B0218+357 at $z_{\\textrm{\\scriptsize em}}=0.94$ with the lensing object at $z_{\\textrm{\\scriptsize abs}}=0.69$ (Wiklind \\& Combes 1995, Menten \\& Reid 1996, Gerin et al. 1997), PKS 1830-211 at $z_{\\textrm{\\scriptsize em}}=2.5$ with the lensing object at $z_{\\textrm{\\scriptsize abs}}=0.89$ (Wiklind \\& Combes 1996b) and PMN J0134-0931 at $z_{\\textrm{\\scriptsize em}}=2.2$ with the lensing object at $z_{\\textrm{\\scriptsize abs}}=0.77$ (Kanekar et al. 2005, 2012). OH lines are detected in all five systems but only two have conjugate OH molecular absorption/emission lines (PKS 1413+135, \\citealt{dar04}, \\citealt{kan04}; PMN J0134-0931, \\citealt{kan12}) important for fundamental constants work \\citep{dar03}.  \\footnotetext{In this paper molecular absorption refers to the multitude of species seen in the radio-frequency spectra of giant molecular clouds, not the UV absorption transitions of H$_2$ and CO absorption arising in cool halo gas probed by far-UV spectroscopy in the Milky Way and high-$z$ QSO spectra of intervening galaxies (\\citealt{not10}, \\citealt{sri08})}  Although ``blind'' searches of samples of strong radio sources have been undertaken to find molecular absorption, mostly they have led to negative results, although a few new \\ion{H}{1} 21~cm absorbers have been found (\\citealt{cur11a}, \\citealt{gup09}). Some of the unsuccessful surveys for high-$z$ molecular absorbers include: millimeter observations of strong sources with known redshifts (Willett, Kanekar, \\& Carilli, private communication), cm- and m-wave observations of reddened quasars and radio galaxies, sources with damped Ly$\\alpha$ absorbers, and type-2 quasars, etc. (Curran et al. 2006, 2011; \\citealt{cur10}). While a few new \\ion{H}{1} 21~cm absorbers have been discovered, no new molecular absorbers have been discovered in these searches. This suggests both that high-z molecular absorbers are very rare and that rather standard techniques cannot be used to discover molecular absorbers (e.g., searches using samples containing damped Ly $\\alpha$ absorbers; \\citealt{gup09}); i.e., the obscuration we are trying to find significantly absorbs and reddens the quasar optical-UV continuum so that it is not detectable in most cases. Thus, optically-bright quasars are not a good sample to use to search for radio-frequency molecular absorption \\citep{cur08, cur11b}.  With the overall goal of discovering new examples of high-$z$ molecular absorbers, we have embarked on an extensive discovery program using a new selection method devised specifically for this task. In this paper we describe the first steps towards the ultimate goal of discovering a substantial sample of high-$z$ radio absorption line systems. This paper is arranged as follows: In Section 2 we describe this new selection method which combines Very Large Array (VLA) Faint Images of the Radio Sky at Twenty centimeters (FIRST; \\citealt{bec95}) radio source detections with Sloan Digital Sky Survey (SDSS; \\citealt{yor00}) photometry and morphologies for the coincident galaxies designed to select strong radio sources which are associated with late-type galaxies. In Section 3 we describe new broad-band near-IR (NIR) imaging of this sample obtained by us at the Apache Point Observatory's 3.5m telescope (APO hereafter) which allows us to search for nuclear point sources which are obscured optically and to extend the spectral energy distributions (SEDs) of these sources from the SDSS optical photometry to longer wavelengths. The analysis of the NIR images and of the optical/NIR SEDs is described in Section 4. A few new spectroscopic redshifts obtained in the progress of this work are also reported in this Section.  Individual results for some representative and also unusual objects are presented in Section 5 and candidates for finding atomic and molecular absorption are identified. We conclude with a summary of the observational results and discuss the future work using our on-going survey in Section 6.  Paper 2 in this series will describe high-resolution VLA and Very Long Baseline Array (VLBA) radio maps of these sources to identify those sources containing very compact structures and to identify sources which have radio/optical-NIR position offsets indicative of foreground, not associated absorption. A first round of \\ion{H}{1} 21~cm absorption observations from this sample will be presented in Paper 3. ",
        "conclusions": "With the goal of finding more high-z molecular absorbers, we have selected 70 strong radio sources with non-elliptical morphology in the optical and obtained NIR images at the APO 3.5m telescope to further characterize their basic nature. Selection from the FIRST survey at $f_{20{\\textrm {\\scriptsize cm}}}>300$ mJy assures that we are targeting radio-loud AGN. An SDSS associated galaxy classification as non-elliptical gives us the greatest probability of finding plentiful gas, regardless of the reason for the late-type morphologies; e.g., the observed galaxy could either be a very young radio galaxy like a CSO or a foreground galaxy like in a gravitational lens system. In this first paper on ``Invisible AGN'' we have presented a selection method for finding highly-obscured AGN and the optical-NIR data necessary both to confirm their morphology and judge the amount of obscuration present.  We have obtained NIR images of the sample to supply us with morphological information to check the SDSS classification, to extend the SDSS optical SEDs into the near-IR and to search for heavily obscured nuclei. Our NIR observations are consistent morphologically and photometrically with the expectation that many of these sources are heavily obscured systems.  Analysis of the NIR morphology and of the optical/NIR SEDs provides a consistent picture of a heavily obscured, radio-loud AGN in most cases (see detailed discussion in Section 4.2). Also we have made a comparison between the photomeric redshifts and redshifts estimated from the $K$-band magnitudes.  Because radio-loud AGN exhibit a significant ``Hubble diagram'' type correlation out to $z\\sim$3, we can use the $K$-band magnitudes as a crude redshift indicator. For nine sources these K-band derived redshifts are significantly greater than their photo-$z$s (see Table 4 and Figure 7), indicating either that the host galaxy is not a giant elliptical or that the SDSS detected galaxy is forground to the AGN. In either case, these are excellent candidates for finding foreground absorption. However, this analysis also makes it clear that it is not advisable to use photometric redshifts of {\\it Q} or {\\it Q$+$abs}-type sources to conduct an efficient \\ion{H}{1} 21~cm absorption-line search. Additionally, Figure 7b shows that even when just the {\\it G}-type sources are used, there is a systematic offset ($\\Delta z \\sim 0.05-0.1$) between the SDSS-derived photometric redshifts and the spectroscopic redshifts on top of the standard statistical errors present in the SDSS photo-$z$ measurement ($\\Delta z \\sim 0.1$--0.12). While it may be possible to conduct an \\ion{H}{1} 21~cm search of sources with {\\it G}-type SEDs using photometric redshifts derived using both SDSS+NIR photometry (called $z_{\\textrm{{\\scriptsize fit}}}$ herein), the statistical errors present (see Figure 7b) are too large to make such a search successful in the presence of the typical RFI environment found at ground-based radio telescopes. Again a more accurate spectroscopic redshift is required.  While the optical/NIR properties described herein are compelling that most of the sources within sample are obscured, 'invisible AGN', it is also important to continue to scrutinize the basic properties of the radio sources against which absorption might be observed. Regardless of the amount of obscuration between us and the nucleus of an AGN, its radio spectrum is unlikely to show strong absorption lines if the radio source contains mostly geometrically-extended flux. This is because we are viewing the large majority of the source flux along unobscured lines-of-sight that avoid going through the host galaxy (and its obscured nucleus) entirely. VLA A-configuration and VLBA maps are in-hand and will provide identification of the most compact radio sources in this sample. These high resolution maps also will be used to make an accurate comparison between the radio source structure and the optical/NIR galaxy positions to identify potential intervening absorption candidates. These maps and their analysis will be presented in Paper 2.  Not only will VLA/VLBA observations discover new examples of CSOs and CSSs but they will also identify those obscured nuclei which are likely to show the strongest absorption due to possessing very compact radio continuum structures.  Hopefully, some of the absorptions we discover will be molecular lines."
    },
    "1208/1208.6044_arXiv.txt": {
        "abstract": "Understanding the origin of AGN absorption line profiles and their diversity could help to explain the physical structure of the accretion flow, and also to assess the impact of accretion on the evolution of the AGN host galaxies. Here we present our first attempt to systematically address the issue of the origin of the complexities observed in absorption profiles. Using a simple method, we compute absorption line profiles against a continuum point source for several simulations of accretion disk winds. We investigate the geometrical, ionization, and dynamical effects on the absorption line shapes. We find that significant complexity and diversity of the absorption line profile shapes can be produced by the non-monotonic distribution of the wind velocity, density, and ionization state. Non-monotonic distributions of such quantities are present even in steady-state, smooth disk winds, and naturally lead to the formation of multiple and detached absorption troughs. These results demonstrate that the part of a wind where an absorption line is formed is not representative of the entire wind. Thus, the information contained in the absorption line is incomplete if not even insufficient to well estimate gross properties of the wind such as the total mass and energy fluxes. In addition, the highly dynamical nature of certain portions of disk winds can have important effects on the estimates of the wind properties. For example, the mass outflow rates can be off up to two orders of magnitude with respect to estimates based on a spherically symmetric, homogeneous, constant velocity wind. ",
        "introduction": "Accretion disk winds are among the most promising physical mechanisms able to link the small and the large scale phenomena in active galactic nuclei (AGN) and to shed light on the physics of mass accretion/ejection around supermassive black holes (SMBHs). Such winds are currently directly observed, as blueshifted and broadened absorption lines in the UV and X-ray spectra of a substantial fraction of AGN \\citep[e.g.,][]{2003ARA&A..41..117C}.  In the UV band we observe, with decreasing width of the absorption troughs, broad absorption lines (BALs) in about 15~\\% of optically selected AGN \\citep[e.g.,][]{2008MNRAS.386.1426K,2009ApJ...692..758G}; mini-broad absorption lines (mini-BALs) and narrow absorption lines (NALs) in about 30~\\% of AGN \\citep{2008ApJ...672..102G}. Once corrected for selection effects, the intrinsic fraction of AGN hosting BALs is estimated to be $\\sim 30-40$~\\% \\citep{2011MNRAS.410..860A}; summing also the corrected fraction for mini-BALs and NALs, the intrinsic fraction of AGN hosting blueshifted UV absorption lines results to be as high as $\\sim 70$~\\% \\citep{2012arXiv1204.3791H}. In all these cases, the absorbers are associated to resonant transitions of moderately ionized elements like Mg~\\textsc{ii}, Al~\\textsc{iii}, C~\\textsc{iv}, Si~\\textsc{v}, and can be outflowing with velocities as high as $\\sim 0.2$~$c$, the main difference among them being in the width of their absorption troughs. Given their large width, BAL structures are the easiest to be identified even in low-resolution UV spectra, and have therefore been studied in much more detail than mini-BAL and NAL structures. The shape of BAL troughs is quite complex, showing a variety of profiles among different sources \\citep[e.g.,][]{1984ApJ...280...51T, 1987ApJ...317..450F, 1988ApJ...325..651T,2002ApJS..141..267H}. When high spectral resolution, high signal-to-noise observations have become available, generally also mini-BALs and NALs showed a variety of profiles \\citep[e.g.,][]{2007ApJ...660..152M,2010ApJ...722..997W,2011MNRAS.410.1957H, 2011MNRAS.411..247R}. Recently a number of AGN have been shown to display strong variability in their UV absorption troughs, including emergence of BAL structures and transitions from mini-BALs to BALs \\citep[e.g.,][]{2002MNRAS.335L..99M,2008MNRAS.391L..39H,2009ApJ...701..176L,2010ApJ...724L.203K,2012MNRAS.421L.107V}.  In the X-ray spectra of $\\sim 50$~\\% of AGN, we observe also lower-velocity (from $\\sim 100$ to $\\sim 1000$~km~s$^{-1}$) outflowing ionized gas (the so-called warm absorber) in the transitions of ionized elements such as C~\\textsc{vi}, O~\\textsc{vii}, O~\\textsc{viii}, Si~\\textsc{vi} \\citep[e.g.,][]{2007MNRAS.379.1359M}. Higher velocity absorbing gas in the transitions of Fe~\\textsc{xxv} and Fe~\\textsc{xxvi}, with blueshift in the range $0.003-0.3$~$c$ is observed in a number of AGN \\citep[e.g.,][]{2003MNRAS.345..705P, 2005ApJ...630L.129R, 2006AN....327.1012C, 2007MNRAS.375..227M, 2008A&A...483..161T,2011A&A...536A..49G}; their fraction among the AGN population is currently unknown, the only reliable estimate being of about $30-40$~\\% for a sample of low redshift AGN \\citep{2010A&A...521A..57T}. X-ray BALs outflowing up to $0.7$~$c$ have been also observed in a handful of AGN \\citep{2002ApJ...579..169C,2003ApJ...595...85C,2009ApJ...706..644C, 2012A&A...544A...2L}. The ionization state of the X-ray absorbers spans a large range of values, from e.g. C~\\textsc{vi} in the case of warm absorbers up to Fe~\\textsc{xxvi} for X-ray BALs. The observed high-velocity X-ray absorbers profiles are quite complex and display strong temporal variability both in depth and velocity shift; also for the case of the lower-velocity X-ray warm absorbers, deep observations have revealed significant complexities in the absorption trough profiles \\citep[e.g.,][]{2005ApJ...627..166R,2011MNRAS.413.1251P} and variability in ionization and velocity shift \\citep[e.g., see][for some recent results]{2010A&A...510A..92L,2011A&A...533A...1M}.  Velocity and ionization state are only two of the physical properties of AGN winds; the mass and energy fluxes carried by them are other fundamental properties that we wish to be able to measure in order to quantify the actual impact of these winds on the surrounding environment, i.e., the amount of feedback. To measure the mass and energy flux, the density of the wind material and the distance between the central SMBH and the place where absorption occurs must be known. Unfortunately, these two quantities are difficult to measure, and estimates for them have always relied on several assumptions. In particular, a geometrically and optically thin, constant-density, spherically symmetric, single-zone of gas in ionization equilibrium and outflowing with a uniform velocity has been usually assumed to derive constraints on the distance of the absorbing material from the central ionizing source.  Using these assumptions, reported estimates of the distances of absorbers from the central SMBH range from the inner regions of the accretion disk for the high-velocity X-ray BALs \\citep[e.g.,][]{2009ApJ...706..644C}, to the parsec-scale torus for the X-ray warm absorbers \\citep[e.g.,][]{2005A&A...431..111B}, to the kiloparsec-scale for some UV BALs \\citep[e.g.,][]{2010ApJ...709..611D}. Currently the uncertainties on the distances are relatively large, which translates to a high uncertainty on the mass outflow rate and on the kinetic energy injection associated with such winds \\citep[e.g.,][]{2003ApJ...593L..65R,2007ApJ...659.1022K,2009A&A...504..401C}.  Theoretical arguments and numerical simulations have been used to show that accretion disks are able to launch and accelerate powerful outflows by various physical mechanisms such as radiation pressure, magnetic pressure, and thermal pressure \\citep[see e.g.][for reviews on the subject]{2006MmSAI..77..598K,2007ASPC..373..267P, 2007Ap&SS.311..269E}. The driving mechanism and the accreting system physical parameters (e.g., accretion rate, black hole mass, magnetic field configuration, UV/X-ray flux ratio) determine the wind characteristics \\citep[e.g.,][]{2005ApJ...630L...9P}. Despite the many parameters that combine themselves in complicated ways, only a few geometries of streamlines are shown by AGN accretion disk wind models.  \\begin{figure} \\figurenum{1} \\epsscale{.9} \\plotone{f1.eps} \\caption{\\label{f1}The three ``typical'' streamlines for accretion disk winds. 1: MHD-typical, polar and convex streamline. 2: hydrodynamical-typical, equatorial and concave streamline. 3: a radial streamline typical of spherically symmetric winds.} \\end{figure}  Figure~\\ref{f1} shows the main general shapes of streamlines that are predicted by models of AGN accretion disk winds. Number 1 is a convex, polar streamline, typical of magnetic driving scenarios \\citep[e.g.,][]{1982MNRAS.199..883B,1991ApJ...379..696L,2003ApJ...595..631K, 2005ApJ...630..945A, 2005ApJ...631..689E}. Number 2 is a concave, more equatorial streamline, typical of line-driven \\citep[e.g.,][]{2000ApJ...543..686P, 2004ApJ...616..688P} and thermally driven disk winds \\citep[e.g.,][]{1996ApJ...461..767W,2004ApJ...607..890F, 2010ApJ...719..515L}. What is usually assumed in order to deduce quantities such as mass outflow rate, ionization state, or absorber distance from the continuum source is the configuration number 3, i.e. a radial flow.  However, the different geometries correspond to very different wind properties: the mass flux $\\dot{m}_{out}$ across the area of a flow tube, $A(r)$ is \\begin{equation}\\dot{m}_{out}= A(r) \\upsilon(r) \\rho(r)\\end{equation} where $\\upsilon$ is the velocity normal to the area, and $\\rho$ is the density. For a simple radial wind with a constant velocity, $A(r) \\propto r^2$ and $\\rho(r)\\propto r^{-2}$, whereas for a cylindrical wind (with constant velocity) $A(r)=const$ and $\\rho(r)=const$. This density dependence on the flow geometry has received little consideration in previous works on estimating the flow mass and energy rates using observations \\citep[however, see][]{2012arXiv1207.7348W}.  In this article we will report on results from our first attempt of systematically addressing the following questions: what drives the appearance of different line profiles? Can we distinguish among different geometries for accretion disk winds from observations? How reliable are estimates of wind properties, when a radial or spherical flow is assumed? The secondary questions that will lead to answer the former ones are then: what are the implications of radial flow assumptions when deriving physical quantities related to the wind? Do different wind physical characteristics produce different line profiles? If yes: how and why? Our method to compute line profiles is introduced in Section 2; results are presented in Section 3 and discussed in Section 4; finally, conclusions are presented in Section 5. ",
        "conclusions": "} We have considered several physical models of accretion disk winds in AGN, and computed absorption line profiles predicted by such models. For the simplest isothermal disk winds, we found that although the models are non-spherical, they predict shapes of the synthetic profiles that could have very similar characteristics for different wind geometries.  Therefore we conclude that the wind geometry is not the main contributor to the great diversity of the observed line shapes.  More complexities and diversities of the profile shapes can be produced by facilitating one of the key properties of AGN disk winds: the non-monotonic distribution of basic wind properties such as velocity, density, and ionization state. Owing to the complex distribution of these quantities along the LOS, multiple and detached absorption troughs can be easily produced in radially extended, non-spherical outflows. The highly dynamical nature of certain portions of AGN disk winds can also have significant effects on the mass and energy fluxes estimates, that can be off up to two orders of magnitude with respect to estimates based on a spherically symmetric, homogeneous, and constant velocity wind."
    },
    "1208/1208.4277_arXiv.txt": {
        "abstract": "The measurements of the 21-cm brightness temperature fluctuations from the neutral hydrogen at the Epoch of Reionization (EoR) should inaugurate the next generation of cosmological observables. In this respect, many works have concentrated on the disambiguation of the cosmological signals from the dominant reionization foregrounds. However, even after perfect foregrounds removal, our ignorance on the background reionization history can significantly affect the cosmological parameter estimation. In particular, the interdependence between the hydrogen ionized fraction, the baryon density and the optical depth to the redshift of observation induce nontrivial degeneracies between the cosmological parameters that have not been considered so far. Using a simple, but consistent reionization model, we revisit their expected constraints for a futuristic giant 21-cm omniscope by using for the first time Markov Chain Monte Carlo (MCMC) methods on multiredshift full sky simulated data. Our results agree well with the usual Fisher matrix analysis on the three-dimensional flat sky power spectrum but only when the above-mentioned degeneracies are kept under control. In the opposite situation, Fisher results can be inaccurate. We show that these conditions can be fulfilled by combining cosmic microwave background measurements with multiple observation redshifts probing the beginning of EoR. This allows a precise reconstruction of the total optical depth, reionization duration and maximal spin temperature. Finally, we discuss the robustness of these results in presence of unresolved ionizing sources. Although most of the standard cosmological parameters remain weakly affected, we find a significant degradation of the background reionization parameter estimation in presence of nuisance ionizing sources. ",
        "introduction": "\\label{sec:intro}  Among the next generation of cosmological probes, interferometric radio telescopes observing the redshifted 21-cm line associated with the hyperfine transitions of neutral hydrogen atoms have attracted a lot of attention (see Refs.~\\cite{Loeb:2000fc, Furlanetto:2006jb, Barkana:2006ep, Pritchard:2011xb} for reviews). These telescopes are revolutionary in their design as in the proposed Fast Fourier Transform Telescope (FFTT) which is conceptually an all digital antennas array imaging the whole visible sky at once~\\cite{Tegmark:2008au}. Images would then be reconstructed by a two-dimensional fast Fourier transform over the $N$ antennas signal, i.e. in $N \\ln N$ operations. As argued in Ref.~\\cite{Tegmark:2009kv}, compared to the required $N^2$ pairing in traditional interferometers, the gain could be used to scale up the telescope size and sensitivity thereby allowing measurements of cosmological signals. The first generation of such telescopes is under deployment~\\cite{Peterson:2006bx, Garrett:2009gp, Lazio:2009bea,Rawlings:2011dd, Pen:2008ut, Mitchell:2010cc, Ord:2010st}. Although not designed for cosmological purposes, they aim at detecting a cosmological signal from the Epoch of reionization~\\cite{Parsons:2009in, Bittner:2010fi, Santos:2010hj, Harker:2010ht, Chang:2010jp, Paciga:2010yy, Morandi:2011hn}. It is still a matter of active research to know if the signal coming from the astrophysical foregrounds can be properly separated from the cosmological one~\\cite{DiMatteo:2001gg, Jelic:2008jg, Bowman:2008mk, Petrovic:2010me, Harker:2011et}. One should indeed keep in mind that the former is actually a few order of magnitude stronger than the latter and this has triggered interest in the cross-correlation of the 21-cm signal with much cleaner data such as the Cosmic Microwave Background (CMB) or galaxy surveys~\\cite{Alvarez:2005sa, Tashiro:2009uj, Jelic:2009sd,Wyithe:2006vg, Furlanetto:2006pg}.  The physical origin of the cosmological 21-cm radiation lies in the differential cooling induced by the expansion of the Universe on relativistic and non-relativistic gases. After recombination, one would naively expect  the temperature of neutral hydrogen $\\Tgas$ to decrease in $1/a^2$, where $a$ is the Friedmann--Lema\\^{\\i}tre--Robertson--Walker (FLRW) scale factor, which is faster than the radiation temperature $\\Trad$ scaling as $1/a$. As a result, the neutral hydrogen gas becomes cool enough to be able to absorb CMB photons by a spin flip hyperfine transition at a wavelength of 21-cm. Tuning a radio telescope at the corresponding redshifted frequency allows, in principle, to probe the density fluctuations of neutral hydrogen over the sky and at any redshift, hence the so-called cosmological tomography~\\cite{1979MNRAS.188..791H, 1990MNRAS.247..510S}. In reality, the situation is a bit more complex and one has to take into account the evolution of the Boltzmann distribution of neutral hydrogen due to the collisions with the residual electrons, protons and absorption versus stimulated emission of CMB photons~\\cite{Madau:1996cs, Seager:1999bc, Seager:1999km, Furlanetto:2007te}. The background evolution can nevertheless be numerically derived and we have plotted in Fig.~\\ref{fig:btbg} the resulting brightness temperature evolution.  \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{btbgdages.eps} \\caption{Background evolution of the neutral hydrogen gas and spin temperatures (top) as a function of the redshift during the dark ages. The spin temperature $\\Tspin$ being defined as $n_1/n_0 = 3 \\exp[-\\hbar \\omega_{21}/(k_B\\Tspin)]$ where $n_0$ and $n_1$ denote the number density of atoms in the singlet and triplet hyperfine states, respectively. The bottom panel shows the resulting brightness temperature.} \\label{fig:btbg} \\end{center} \\end{figure}  For an assumed generic set of cosmological parameters, taken from the Wilkinson Microwave Anisotropy Probe (WMAP) seven-year data~\\cite{Komatsu:2010fb}, the absorption is maximal around a redshift of $z\\simeq 10^2$. Notice that at lower redshifts, the spin temperature is driven again towards the photon temperature and the signal vanishes till the EoR. The evolution of neutral hydrogen density fluctuations can similarly been predicted during the dark ages by using the theory of cosmological perturbations~\\cite{Loeb:2003ya, Bharadwaj:2004nr, Naoz:2005pd}. Provided the length scales of interest remain in the linear regime, the theoretical predictions for the 21-cm power spectra are neat~\\cite{Hirata:2006bn, Lewis:2007kz, Lewis:2007zh, Mao:2008ug}. Being three-dimensional in nature, the information content is huge and have been used to forecast constraints on various cosmological models such as non-Gaussianities~\\cite{Pillepich:2006fj, Cooray:2006km, Joudaki:2011sv,Chongchitnan:2012we}, cosmic strings~\\cite{Khatri:2008zw, Brandenberger:2010hn, Berndsen:2010xc, Pagano:2012cx}, dark matter signatures~\\cite{Shchekinov:2006eb, Valdes:2007cu, Borriello:2008gy, Cumberbatch:2008rh, Natarajan:2009bm}, modified gravity~\\cite{Brax:2012cr} and inflation~\\cite{Barger:2008ii, Gordon:2009wx, Masui:2010cz, Adshead:2010mc}.  As already mentioned, in addition to having a small amplitude compared to foregrounds, the dark ages signal is redshifted to wavelengths of typically twenty meters and it is not obvious that it may actually be used for cosmology in a foreseeable future. One very futuristic approach would be to build those radio telescopes on the Moon~\\cite{Jester:2009dw}. A more reasonable approach is the possibility to probe the hydrogen density fluctuations with the 21-cm line at the EoR~\\cite{Ciardi:2003hg, Sethi:2005gv, Barkana:2005xu}. This time, the exciting photons are coming from reionization sources (stars or quasars) instead of the CMB and are expected to produce an efficient spin states population inversion. This results into $\\Tspin \\gg \\Trad$ and generates an emission line at a much smaller redshifted wavelength. We have plotted the expected evolution of the brightness temperature during the EoR in Fig.~\\ref{fig:btbgeor} for a simple reionization model that we describe in Sec.~\\ref{sec:eormodel}.  \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{btbgeor.eps} \\caption{Reionization driven spin temperature (top) and brightness temperature (bottom) with respect to the observation redshift. The reionization model is described in Sec.~\\ref{sec:eormodel} ($\\tau = 0.088$).} \\label{fig:btbgeor} \\end{center} \\end{figure}  The resulting signal is around $25\\,\\mK$ in amplitude, which is of the same order than those coming from the dark ages, but at a higher frequency thereby rendering its detection more likely. However, the theoretical predictions are now far less neat due to the additional astrophysical uncertainties associated with the way reionization proceeds~\\cite{2011MNRAS.410.1377T, Mao:2011xp}. However since $21\\,\\cm$ data are three-dimensional by nature, it has been shown in Refs.~\\cite{Liu:2011hh, Liu:2011ih} that the redshift evolution could be used to efficiently eliminate the expected foregrounds while still keeping some of the so-called longitudinal modes for cosmology.  Forecasting the 21-cm constraints for cosmology is usually made from Fisher matrix analysis which merely assumes the likelihood to be Gaussian around the best fit. Provided the parameters are well constrained, the method is fast and accurate~\\cite{Tegmark:1996bz}. However, as noted in Ref.~\\cite{Mao:2008ug}, some model parameters linked to reionization, such as the optical depth or the hydrogen ionized fraction, can be completely degenerated from the 21-cm point of view and factorize out of the Fisher matrix. In a realistic situation, this should not be the case as all reionization parameters are uniquely determined by the background reionization history. One may therefore wonder how their correlations to the standard cosmological parameters affect the forecasts.  In Refs.~\\cite{Pritchard:2010pa, Morandi:2011hn}, the problem of reconstructing the background reionization history is specifically addressed in the context of future 21-cm experiments that would be devoted to the global signal, i.e. the homogeneous mode. Here, we would like to discuss this issue for the 21-cm FFTT-like experiments which are poorly, if not at all, sensitive to the constant mode. In that situation, the cosmological brightness fluctuations $\\dTb(\\bx,z)$ have simultaneously a dependence in both the background evolution and spectral shape. This question is peculiar to 21-cm observables because the homogeneous mode $\\Tb(z)$ depends on various cosmological parameters, and when concerned with EoR, on the reionization (see Sec.\\ref{sec:eorsignal}). Let us notice that such a situation does not occur for CMB not only because $\\Tcmb$ is measured, but because it is uniquely determined by $\\OmegaR$.  In this context, we would like to quantify how cosmological forecasts are affected by the background reionization history, and if it is possible for omniscopes to constrain the reionization parameters. For this purpose, we consider a giant FFTT-like ground-based instrument and assume that most of the foregrounds can indeed be eliminated using redshift evolution. As a result, we will be keeping only a few redshift slices for our forecasts, eventually marginalizing over some nuisance ionization spectra modelled as in Ref.~\\cite{Mao:2008ug}. We consider a simple reionization model for which the time evolution of the spin temperature, hydrogen ionized fraction and optical depth are given as a function of some extra parameters. In this manner, it is imposed that they are intrinsically correlated by reionization physics. As already mentioned, in order to go beyond Fisher matrix, we use MCMC methods over a modified version of the $\\CAMB$ 21-cm code~\\cite{Lewis:1999bs, Lewis:2007kz, Lewis:2007zh} which includes the Epoch of Reionization as well as some nuisance ionization spectra. This allows us to use and compare both Fisher matrix analysis and full sky simulated data coupled to MCMC methods~\\cite{Lewis:2002ah}.  The paper is organized as follows. In Sec.~\\ref{sec:models}, we present the reionization and telescope models used in the analysis while the methodology and results are presented in Sec.~\\ref{sec:reiocast}. Starting from one redshift slice only, we show that the reionization induced degeneracies damage, and sometime prevent, the determination of some cosmological parameters, such as the baryon density. In these situations, we also find that a Fisher analysis can be inaccurate for some strongly correlated parameters. Combining a few redshifts, which sample the beginning of reionization, improves the situation but does not compete with a Planck-like CMB experiment. However, we show in Sec.~\\ref{sec:addcmb}, that combining 21-cm and CMB lifts almost all those degeneracies thereby improving the overall accuracy by an order of magnitude. In that situation our MCMC results match with the usual Fisher matrix analysis while the underlying reionization model can be completely determined. Finally, we show that these results may be mitigated by the presence of toy nuisance reionization spectra. ",
        "conclusions": "In this paper, we have quantitatively shown that omniscopes could be used to constrain the background reionization history while being not sensitive to the zero mode of the brightness temperature.  For this purpose, we have considered a simple but consistent reionization model completely determined by the total optical depth $\\tau$, the reionization duration $\\Deltaz$ and the asymptotic spin temperature $\\TspinMax$. We have used both Fisher matrix approaches on the three-dimensional power spectrum and MCMC methods on full sky simulated data to forecast the expected variance of all model parameters. Our results suggest that it is crucial to combine multiple redshifts at the EoR with CMB data to keep all degeneracies under control, in particular owing to an accurate determination of $\\OmegaB h^2$. When this is not the case, we have shown that the Fisher predictions can be quite inaccurate, eventually being overoptimistic, but also overpessimistic when the conditions for the Cram\\'er-Rao inequalities are no longer met. For those situations, only the MCMC methods ends up being usable.  Within combined data, the perspectives are quite good as the background history can be fully reconstructed, even in the presence of correlated nuisance ionizing sources. Although we have discussed only one kind of FFTT telescope, our results could be easily generalized to other configurations by a proper rescaling of the resolution and noise parameters of Sec.~\\ref{sec:models}. Still, our work would need to be extended with a more complete reionization model, eventually adjusted to numerical simulations in order to include a realistic dependence between the reionization parameters and the redshift evolution of the ionization power spectra. Another extension would be the inclusion of exotic reionization sources, such as a decaying or annihilating dark matter component."
    },
    "1208/1208.2066_arXiv.txt": {
        "abstract": "We report extensive new photometry and spectroscopy of the highly variable young stellar object PTF 10nvg (also known as IRAS 20496+4354 and V2492 Cyg), including optical and near-infrared time series data as well as mid-infrared and millimeter data. Following the previously reported 2010 rise to \\rptf $\\lapprox 13.5^m$ and subsequent fade, during 2011 and 2012 the source underwent additional episodes of brightening, followed by several magnitude dimming events including prolonged faint states at \\rptf $\\gapprox 20^m$.  The observed high-amplitude variations are largely consistent with extinction changes ($\\Delta A_V$ up to 30 mag) having a $\\sim 220$ day quasi-periodic signal.  However, photometry measured when the source was near maximum brightness in mid-2010 as well as in late-2012 does not phase well to this period. Spectral evolution includes not only changes in the spectral slope but correlated variation in the prominence of TiO/VO/CO bands and atomic line emission, as well as anticorrelated variation in forbidden line emission which, along with H$_2$, dominates optical and infrared spectra at faint epochs. Notably, night-to-night variations in several forbidden doublet strengths and ratios are observed. High-dispersion spectra were obtained in a variety of photometric states and reveal time-variable line profiles. Neutral and singly-ionized atomic species are likely formed in an accretion flow and/or impact while the origin of zero-velocity atomic \\ion{Li}{1} $\\lambda$ 6707 in {\\it emission} is unknown. Forbidden lines, including several rare species, exhibit blueshifted emission profiles and likely arise from an outflow/jet. Several of these lines are also seen spatially offset from the continuum source position, presumably in a shocked region of an extended jet.  Blueshifted absorption components of the \\ion{Na}{1} D doublet, \\ion{K}{1} $\\lambda$ 7665,7669 doublet, and the \\ion{O}{1} 7774 triplet, as well as blueshifted absorption components seen against the broad H$\\alpha$ and \\ion{Ca}{2} triplet emission lines, similarly are formed in the outflow.  CARMA maps resolve on larger scales a spatially extended outflow in mm-wavelength CO. We attribute the recently observed photometric and spectroscopic behavior to rotating circumstellar disk material located at separation $a \\approx 0.7 ({M_\\ast}/{M_\\odot})^{1/3}$ AU from the continuum source, causing the semi-periodic dimming. Occultation of the central star as well as the bright inner disk and the accretion/outflow zones renders shocked gas in the inner part of the jet amenable to observation at the faint epochs. We discuss PTF~10nvg as a source exhibiting both accretion-driven (perhaps analogous to V1647 Ori) and extinction-driven (perhaps analogous to UX Ori or GM Cep) high-amplitude variability phenomena. ",
        "introduction": "PTF~10nvg, in the North America Nebula region of recent star formation, was identified by the Palomar Transient Factory \\citep[PTF;][]{Law09,Rau09} on 2010 July 8, as an optical transient based on automatic discovery and classification codes \\citep{Bloom11}. The source was rapidly followed up by the PTF collaboration. We announced our findings concerning the 2010 outburst of this Class I type young star in \\citet{Covey11} where we presented optical and infrared light curves and multi-epoch optical and infrared spectroscopy.  Therein, we made the analogy of PTF~10nvg (now called V2492 Cyg) to the behavior of V1647 Ori, an embedded Class I-type young star that was observed over the past decade to undergo several large amplitude photometric events on few year time scales.  V1647 Ori and PTF~10nvg both displayed $\\sim$4--6 mag photometric rises from their faint states that were similar in the early stages to the outbursts of FU Ori stars.  However, these particular objects and others like them do not have the spectroscopic characteristics of FU Ori stars; rather than being absorption line dominated (especially at high--dispersion), they are emission-line dominated, with absorption seen in only a handful of blueshifted features arising in strong winds. Further, the V1647 Ori-type objects do not remain in the elevated photometric state for the long time scale associated with FU Ori outbursts (estimated at roughly a century); instead, their brightening episodes last only a few months to a few years and are characterized by large amplitude fluctuations on month to few month time scales.  Members of this category possibly undergo repeated episodes of their outbursting (and/or extinction dominated) behavior at several year intervals. In this regard they are similar to the lower amplitude ($\\sim$2-4 mag) but repeating outbursts of EX Lup-type systems, which last months to $>$1 year each and repeat on few year to decade intervals.  Whether small-scale, low-amplitude events such as EX Lup-type outbursts, larger scale, high--amplitude events such as the V1647 Ori-type events, or similarly large amplitude but also longer duration FU Ori-type events, the star/disk system in such outbursts is interpreted as undergoing an episode of enhanced mass accretion and associated mass outflow. The accretion mechanisms for the different categories of objects are likely related, and are attributed to instabilities in the inner accretion disks, possibly associated with cyclic magnetosphere-disk interactions.  For PTF~10nvg in particular, \\citet{Covey11} estimated a mass accretion rate of $2.5\\times 10^{-7}$\\msun\\pyr\\ in the elevated state of 2010\\footnote{The accretion rates quoted in \\citet{Aspin11} are a factor of 2.5--10 higher.}, similar to the value estimated for EX Lup during its 2008 outburst (e.g. \\citep{Juhasz12}).  Time-variable accretion, however, may not be the entire explanation for many large-amplitude young star variables. As demonstrated herein, time-variable extinction is clearly an important part of the PTF~10nvg interpretation, and may also play a significant role in the observed photometric behavior of many of the so-called ``outburst\" light curves highlighted to date in the literature.  For example, GM Cep was discussed by \\citet{SiciliaAguilar08} as an EX Lup-type object, but later assessed by \\citet{Xiao10}, \\citet{SemkovPeneva12}, and \\citet{Chen12} as having extinction-dominated rather than accretion-dominated time series behavior.  Both phenomena are perhaps simultaneously relevant, as we argue here for PTF~10nvg. Other well-known low-amplitude young star variables such as UX Ori, RR Tau, VV Ser, AA Tau as well as the ``dippers\" discussed by \\citet{Morales11} and \\citet{CodyHillenbrand10,CodyHillenbrand11} also appear to be undergoing short time scale extinction events. Larger amplitude and -- importantly -- long period examples include KH15D \\citep{Hamilton12}, WL4 \\citep{Plavchan08}, and YLW16A \\citep{Plavchan10} which have found explanation in binary interactions with circumstellar disk material.  The 2010 brightening of PTF~10nvg was independently detected by  K. Itagaki and reported in 2010, August in the subscription service of the Central Bureau for Astronomical Telegrams as CBET 2426.  Additional papers to date discussing the post-outburst source include those by \\citet{Kospal11} and \\citet{Aspin11} who also present multi-color photometric, spectroscopic, and SED analysis of the source, often using the \\citet{Covey11} data.  In this paper we present new data gathered by us since the publication of \\citet{Covey11}. Intensive time series photometry shows that the source has continued its large amplitude and color photometric fluctuations.  Time series spectroscopy demonstrates 1) continuum changes; 2) variation in the broad TiO / VO optical band emission that was detected for the first time during the 2010 outburst of PTF~10nvg, along with variation in other molecules such as CO and H$_2$; and 3) drastic changes in the atomic emission lines indicative of accretion and outflow that are correlated (permitted lines) and anti-correlated (forbidden lines) with the photometric brightness. We also present high--dispersion spectral data; the line profiles are used to quantify the velocities relevant to the inflowing and outflowing material. An updated post-outburst spectral energy distribution is discussed, demonstrating that the source brightened relative to its historical spectral energy distribution in the mid-infrared, as well as in the near-infrared and optical, during 2010. A spatially extended, low velocity molecular outflow is detected at millimeter wavelengths.  New high spatial resolution near-infrared imaging rules out the presence of stellar companions within several hundred AU.  In $\\S$\\ref{sec:obs} we describe our observations from 2009-2012. We then present our analysis of the multi-wavelength light curves ($\\S$\\ref{sec:phot}), changes in the overall spectral energy distribution ($\\S$\\ref{sec:sed}), and analysis of the continuum, absorption line, and emission line spectroscopy including discussion of spatially offset forbidden line emission ($\\S$\\ref{sec:spec}). In $\\S$\\ref{sec:mult} we present constraints on the source multiplicity from high angular resolution direct imaging. In $\\S$\\ref{sec:disc} we interpret our results and discuss the broader implications of our findings for accretion and extinction evolution of young stars. Finally in $\\S$\\ref{sec:summ} we summarize and conclude. ",
        "conclusions": "\\label{sec:summ}  1. Continued photometric monitoring of PTF~10nvg by the P48 and PAIRITEL telescopes throughout 2011 and 2012 detected magnitude variations of $\\Delta$R $\\sim$ 10 mag and $\\Delta$K $\\sim$ 3 mag, and color variations of $\\Delta$J-K $\\sim$ 3 mag.  The time series behavior in color-color and color-magnitude diagrams indicates that these photometric variations are consistent with those expected due to changes in line-of-sight extinction amounting to $\\Delta$A$_V >$ 30 mag.  2. Over the long term, attributing PTF~10nvg's non-detection in 2MASS to similar extinction changes suggests the source has historically experienced a range of $\\Delta$A$_V \\sim$ 70 mag.  Alternately, the non-detection in 2MASS may reflect a true outburst scenario for the 2010 brightening and subsequent photometric evolution.  3.  Time scales of several weeks can be associated with dramatic flare-like photometric changes and $\\sim$7 months with the repeated maxima observed in the time series over the past several years. While the initial 2010 maxima were rounded/smooth, later maxima in 2011 and early 2012 were characterized by narrow/sharp photometric peaks. A late-2012 maximum was again broad and rounded/smooth.  4.  The Lomb-Scargle periodogram calculated from JHK$_s$ light curves between 2010 and mid-2012 features a prominent, statistically significant peak at $\\sim$221 days. Under the assumption that the period is the Keplerian orbital timescale, the dust obscuration governing PTF~10nvg's photometric behavior is located at about $0.7 ({M_\\ast}/{M_\\odot})^{1/3}$ AU from the central star. Continued monitoring will be essential to confirm the reality of the periodic signal; the existing light curves span $<$1000 days in total, and include several 100$+$ day gaps.  Further, the late-2012 lightcurve appears to depart from the previously derived periodic trend and may reflect renewed dominance of accretion over extinction effects.  5. WISE detected PTF~10nvg over multiple epochs in 2010. Comparison of these mid-infrared detections with prior detections at similar wavelengths by IRAS, Spitzer, and MSX indicate that PTF~10nvg brightened in the mid-infrared by factors of a few during its summer 2010 brightening.  6. CARMA observations reveal an extended millimeter continuum source having a total (dust plus gas) mass of 0.06 M$_\\odot$.  CARMA also detects spatially unresolved emission in both $^{12}$CO and $^{13}$CO centered on PTF~10nvg's near-infrared position and systemic velocity. Spatially extended, redshifted $^{12}$CO emission is detected to the south of PTF~10nvg, which we interpret as revealing the presence of an outflow cavity.  Over the wider field, the molecular emission is well-aligned with the H$\\alpha$ emission arc.  7. The bright-state optical and near-infrared spectrum of PTF~10nvg is that of a ``continuum-plus-emission\" object, similar to V1331~Cyg but with significant molecular emission contributions to the continuum. The faint-state spectrum is that of a Herbig-Haro object. Both neutral and singly ionized line species probing a range of densities and temperatures are observed.  The lines that characterize PTF~10nvg's faint-state spectrum are present at all epochs, however, suggesting that the spectral evolution is primarily due to the supression of the bright state spectral features during epochs of enhanced extinction.  8. Spectral monitoring reveals that the strong wind signatures detected in the 2010 outburst persisted through 2011 and 2012.  The depth and terminal velocities of optical wind absorption features (i.e., the Na I D and K I 7665/7699 doublets, the O I 7774 triplet) appeared consistent -- or perhaps slightly increased over time in some lines -- across all epochs where the optical continuum was sufficiently strong to enable these features to be measured. In the near-infrared, \\ion{He}{1} $\\lambda$ 10830 absorption increased in both depth and equivalent width by nearly a factor of two from 2010 through 2011.  A strong redshifted emission component to \\ion{He}{1} $\\lambda$ 10830 likely also traces the outflow rate of the inner wind.  9. The evolution of the strengths and profiles of permitted atomic emission lines reveal the presence of time variability in PTF~10nvg's line-of-sight accretion and outflow activity.  The source exhibits optically thick Ca II triplet emission with line ratios that are unusual even for young stars.  The H${\\alpha}$ emission also demonstrates significant variations in both overall strength and kinematic profile, with the line weakening and shifting blueward as the source fades.  Strong redshifted H${\\alpha}$ emission disappeared in June 2011 as PTF~10nvg approached its photometric minimum, while the blueshifted H${\\alpha}$ emission remained distinctly visible.  These observations suggest that PTF~10nvg's \\ion{Ca}{2} and \\ion{H}{1} lines include contributions from both accretion and outflow activity. While the accretion component may dominate during the brightest states, during fainter states only the outflow signature remains. Notably absent from the spectra at any epoch is the \\ion{He}{1} $\\lambda$ 5876 line, a high temperature (T $>$ 15,000 K) tracer seen near-ubiquitously in classical T Tauri stars.  10. Forbidden line emission measuring $\\sim$5000-20,0000 K gas reveals the presence of a likely Herbig-Haro jet with a wide range of densities, $n\\sim10^3-10^7$ cm$^{-3}$. Some line ratios such as [\\ion{O}{1}] 6300 \\AA\\ : 6363 \\AA\\ and [\\ion{Fe}{2}] 7155 \\AA\\ : 7172 \\AA\\ show extreme and unusually time-dependent ratios that are not easily explained in the standard optically thin assumption for forbidden line emission.  Instead, these lines may be partially optically thick.  10. Spatially extended blueshifted emission is visible in several forbidden line species as well as H$\\alpha$ in two-dimensional HIRES images.  These position-velocity diagrams provide further evidence that PTF~10nvg is driving a jet with a line-of-sight velocity of about $-200$ km/s and exciting a strong shock region $\\sim$4\" projected distance from the source (2000AU at 520 pc).  11. The moderate resolution near-infrared spectra obtained in June/July 2011, when PTF~10nvg was near its faintest state, provide an opportunity to study the properties of the jet and extended outflow. PTF~10nvg's H$_2$ line strengths strongly resemble theoretical predictions for C-type shocks and pure X-ray photo-excitation, as well as empirical observations of H$_2$ emission from RW Aur and Knot B of the HH54 protostellar jet; the same is true for the near-infrared [\\ion{Fe}{2}] emission.  12. From near-infrared adaptive optics observations we can exclude the presence of stellar companions within several hundred AU of PTF~10nvg.  Overall, we conclude that the recent time series data on PTF~10nvg will not be easy to interpret without a good deal of hindsight.  Further observational and theoretical study of this enigmatic source is warranted.  \\facility{Facilities: Palomar 48\", PAIRITEL, Lick 3m, Palomar 200\" (two instruments), APO, Keck (four instruments), IRTF, CARMA, DSS, KPNO 0.9m, USNO, 2MASS, UKIDSS, Spitzer, WISE}"
    },
    "1208/1208.0580_arXiv.txt": {
        "abstract": "{ We compute the scalar gravitational radiation from a binary pulsar system in the simplest model that exhibits the Vainshtein mechanism. The mechanism is successful in screening the effect from scalar fields conformally coupled to matter, although  gravitational radiation is less suppressed relative to its general relativity predictions than static fifth forces effects within the pulsar system. This is due to a combination of two effects: firstly the existence of monopole and dipole radiation; secondly the Vainshtein suppression comes from the hierarchy of scales between the inverse frequency scale and the Vainshtein radius, rather than the orbital radius of the pulsar system. Extensions of these results will have direct relevance to infrared modifications of gravity, such as massive gravity theories, which are known to exhibit a Vainshtein mechanism. Generalization to Galileon models with higher order interactions are likely to provide stronger constraints. } ",
        "introduction": "The discovery of cosmic acceleration has spurred a search for consistent modifications of gravity in the infrared, \\cite{Dvali:2007kt}.  Theories in which the graviton acquires a mass, either softly as in the Dvali-Gabadadze-Porrati (DGP) model \\cite{DGP} or cascading gravity \\cite{deRham:2007rw}, or a hard mass as in the newly developed ghost-free models of massive gravity \\cite{dRGT} or their bigravity \\cite{Hassan:2011zd} extensions, are a promising class of such modifications.  These models commonly include light scalar degrees of freedom which arise from additional graviton helicity states. At the linear level, the scalar modes do not decouple in the massless limit, which is known as the van Dan, Veltman, Zakharov (vDVZ) discontinuity, \\cite{van Dam:1970vg}. Despite this, massive gravity models satisfy standard tests of gravity since the new scalars become strongly coupled near dense sources, which suppresses scalar gradients, and yields a force that is much smaller than the Newtonian one. This strong coupling effect is known as the Vainshtein mechanism, \\cite{Vainshtein}. Far from compact sources, the scalars are weakly coupled and massive gravity theories make novel predictions that differ from Newtonian gravity in interesting ways.  This suggests that IR modifications of gravity have extremely small effects on any but cosmological scales.  Much of the phenomenology of the Vainshtein mechanism can be captured by considering the Galileon models \\cite{Galileon}. Indeed the original Galileon model came from considering the decoupling limit of the DGP model \\cite{lpr}, and the structure of this scalar was generalized in \\cite{Galileon} to include all interactions consistent with the Galilean symmetry $\\pi \\rightarrow \\pi + c + v_{\\mu}x^{\\mu}$, having a well-defined Cauchy problem. It was subsequently shown that the generic Galileon arises as the decoupling limit \\cite{deRham:2010ik,deRham:2010gu} of the generic ghost-free massive gravity theory \\cite{dRGT}. In an independent line of reasoning, the Galileon models have been viewed as scalar theories in their own right (independent of their graviton helicity-zero origin) and the most general covariant Galileon theories have been constructed \\cite{CovariantGalileon,Deffayet:2009mn,deRham:2010eu,Goon:2011qf}. In the decoupling limit in which the helicity-zero graviton mode is viewed as weak, all of these theories take on a similar structure and at least in some cases, the physics of the Vainshtein mechanism is qualitatively similar. For this reason, in this article we shall focus on the simplest cubic Galileon model, and will leave generalizations to subsequent work \\cite{PulsarTheReturn}. See Refs.~\\cite{Wyman:2011mp,Koyama:2011xz,Chkareuli:2011te,Sjors:2011iv} for studies investigating the Vainshtein mechanism directly in massive gravity and Refs.~\\cite{Babichev:2009jt} for Fierz-Pauli massive gravity. For other potentially related observations of the Vainshtein mechanism, see Ref.~\\cite{Hui:2012jb,Iorio:2012pv,Babichev:2011iz}.  In the simplest cubic Galileon (DGP) model, the predicted anomalous acceleration $\\Delta a_{\\rm DGP}$ is $\\Delta a_{\\rm DGP}/a_N  \\sim 10^{-15}$, where $a_N$ is the Newtonian acceleration, for a typical binary pulsar system of characteristic mass $3 M_{\\astrosun}$ and semi-major axis $a \\sim 10^{-2}$ AU. We show that this ratio does not set the anomalous contribution to radiated power, which is a few orders of magnitude larger in known binary pulsar systems and could potentially be even stronger in slow, high eccentricity systems.  The enhancement is due to the fact that the relevant length scale is set by the onset of the ``far-field\" region at radius $\\Omega_P^{-1}$ where $\\Omega_P = 2\\pi/T_P$ and $T_P$ the orbital period (often denoted $P_b$ in the literature.)  For $T_P =$ 8 hours, this is $\\sim 10^3$ greater than the orbital radius. Furthermore, the system also radiates in the monopole and dipole channels since we are now radiating into a scalar.  Monopole and dipole radiation exists because the scalar effectively violates the equivalence principle. These effects could potentially get more pronounced in future observations.  Nevertheless, in the simplest model exhibiting the Vainshtein mechanism we consider in this paper, these enhancement are not sufficient to produce scalar gravitationally radiation effects which are within present observational limits from pulsar timing, and the Vainshtein mechanism, albeit slightly more subtle in this time-dependent system, is still very much alive. In a subsequent work we will show how higher order interactions in these different massive gravity/Galileon models are potentially much more strongly constrained by current observations \\cite{PulsarTheReturn}.  The rest of this paper is organized as follows: In section \\ref{sec:Formalism} we review the  cubic Galileon model and derive the general formalism to compute the radiated power in all generality in any multipole, taking care of the Vainshtein mechanism. The power is computed in two independent and equivalent ways using first an effective field theory approach and second a more conventional energy flux derivation. We then apply this formalism in the subsequent sections \\ref{sec:Monopole}, \\ref{sec:Dipole} and \\ref{sec:quadrupole} to compute the respective monopole, dipole and quadrupole radiation in the cubic Galileon model. Finally we compare these results with observations in section \\ref{sec:observations} and summarize our results. ",
        "conclusions": "\\label{sec:observations}  Pulsars in DNS binaries are  used to measure gravitational-radiation effects.  Timing measurements of these systems are  free from contamination due to tidal effects or accretion from a stellar companion. The archetypal such system is the Taylor-Hulse pulsar 1913+16 discovered in 1974,  \\cite{Hulse:1974eb,Taylor:1989sw,Weisberg:2004hi}.  We have summarized the orbital parameters of four known DNS pulsars (A to D) and one pulsar-white dwarf binary (E) in Table 1, \\cite{Lorimer:2005bw},   all of which have measured orbital period derivatives $\\dot T_P$, which agree with GR.  The orbital period derivatives $\\dot T_P$ are given in terms of the non-relativistic energy $E_{\\rm NR}$ and the power emitted by the relation \\ba \\label{eq:TPdot} \\dot T_P =\\frac{3}{2} T_P \\frac{1}{E_{\\rm NR} }\\frac{\\d E_{\\rm NR}}{\\d t}\\,, \\ea where the system's non-relativistic energy is \\ba E_{\\rm NR}=\\frac{1}{2 (8\\pi)^{2/3}}\\frac{M_1 M_2}{\\mpl} \\(\\frac{\\Op^2}{M \\mpl}\\)^{1/3}\\,. \\ea The monopole and quadrupole Galileon radiation are close, but in the explicit examples presented here, the Galileon quadrupole always give the largest contribution. In summary, in the simplest model, the Galileon radiation is at least 7 orders of magnitude below that of GR (7 orders of magnitude for pulsar C and E and 8 for the other DNS pulsars),  which 6 orders of magnitude below the current precision, \\cite{Footnote}, when considering a parameter $m=1.54 \\times 10^{-33}$eV, or equivalently for a strong coupling scale $\\Lambda\\sim 10^{-13}{\\rm eV}\\sim (1000 {\\rm km})^{-1}$. The best precision is still the Hulse-Taylor pulsar and pulsars C and E. Even though the double pulsar D has a good precision $\\sigma$, its low eccentricity make it not the best candidate to probe the Vainshtein mechanism.  As can be seen from the expression \\eqref{eq:TPdot} for the orbital period derivative, using the power emitted in the monopole, \\eqref{e:TheAnswer}, the dipole \\eqref{e:PowerDipole} or the quadrupole, \\eqref{e:PowerQuadrupole}, $\\dot T_P$ scales directly as $\\rs^{-3/2}\\sim m \\sim \\Lambda^{3/2}$. So in order for the scalar field $\\pi$ to have an effect at all on current or upcoming binary pulsar timing observations, the parameter $m$ should be enhanced by roughly 6 orders of magnitudes. Binary Pulsar timing thus put a rough bound of $m < 10^{-27}$eV or $\\Lambda < 10^{-9}$eV. Compared to solar system tests \\cite{solarsystem}, these bounds are not competitive for the cubic Galileon interactions, but could be enhanced when considering higher interactions, \\cite{PulsarTheReturn}.   Present constraints are limited by the sample of DNS pulsar systems. Ideal DNS binary pulsars have long periods, high eccentricities, and are located nearby so that the kinematic corrections are reduced. Over time, measurements of the orbital period derivative become more precise as more data is gathered. The ultimate precision is limited by the uncertainty in the relative acceleration between the sun and the pulsar system, which must be included to obtain the intrinsic orbital period derivative from the apparent one.  We should note that whilst previous authors have considered constraints on the mass of the graviton in binary pulsars, in particular see \\cite{Finn:2001qi}, these authors do not take account of the Vainshtein mechanism, working only in the linearized Fierz-Pauli theory, and they further utilize an incorrect expression for the stress energy of gravitational radiation which accounts only for the helicity two component. As such the constraints obtained there are not appropriate for consistent Lorentz invariant theories of massive gravity as considered in \\cite{dRGT}. They may however be relevant to theories in which there is no propagating helicity zero mode.  In this paper, we have shown the successful implementation of the Vainshtein mechanism in a fully time-dependent setup. This explains how a conformally coupled scalar field, or the helicity-0 mode of a massive graviton can evade typical tests of GR and particularly the well constrained orbital period decay in binary pulsar systems. Nevertheless, despite the presence of an active Vainshtein mechanism, we show that the suppression is less than naively anticipated from purely static systems for several reasons: \\begin{itemize} \\item The suppression factor that arises due to Vainshtein effect is going like $(\\Op \\rs)^{-3/2}$ for the monopole srather than $(\\bar r/\\rs)^{3/2}$ as is the case in static spherically symmetric configurations (for the quadrupole, the suppression factor even acquires an additional velocity suppression). \\item Furthermore, even though the leading contributions from the monopole and dipole radiation vanish from energy and angular momentum conservation, the sub-leading (relativistic) contributions are non-negligible, and in the case of the monopole can be comparable to the Galileon quadrupole radiation. \\item Finally, we have focused here on the simplest realization of the Vainshtein mechanism, namely within the context of the cubic Galileon model. However higher order Galileon interactions could potentially lead to an additional enhancement of the Galileon radiation and ought to be studied in their own right, \\cite{PulsarTheReturn}. \\end{itemize}"
    },
    "1208/1208.4940.txt": {
        "abstract": "{This paper introduces a series of papers aiming to study the dozens of low mass eclipsing binaries (EBLM), with F, G, K primaries, that have been discovered in the course of the WASP survey. %; more precisely this particular paper is the first of a smaller series which will focus on their orbital parameters, long term radial velocity variations and the first estimates of their masses. Our objects are mostly single-line binaries whose eclipses have been detected by WASP and were initially followed up as potential planetary transit candidates. These have bright primaries, which facilitates spectroscopic observations %These objects form a natural extension to the transiting gas giant exoplanet sample found with WASP, and are important for understanding and distinguishing formation mechanisms for very low mass stars, brown dwarfs, and planets.   In addition, they can be used to constrain the mass-radius relation at the bottom of the main sequence and into the brown dwarf regime, where little empirical data currently exists. %By observing in spectroscopy during transit  %hot Jupiters have been found on orbital planes on a variety of angles with respect to the stellar spin. Using the same observing techniques and allows the study of the spin-orbit distribution of F, G, K+M eclipsing binaries through  the Rossiter--McLaughlin effect. % and compare it to hot Jupiter systems. \\\\Here we report on the spin-orbit angle of WASP-30b, a transiting brown dwarf, and improve its orbital parameters. We also present the mass, radius, spin-orbit angle and orbital parameters of a new eclipsing binary, J1219--39b (1SWAPJ121921.03--395125.6, TYC 7760-484-1), which, with a mass of $95\\pm2\\,M_\\mathrm{jup}$, is close to the limit between brown dwarfs and stars. We find that both objects orbit in planes that appear aligned with their primaries' equatorial planes.  Neither primaries are synchronous. J1219--39b has a modestly eccentric orbit and is in agreement with the theoretical mass--radius relationship, whereas WASP-30b lies above it. %We also find that the stellar rotation velocity obtained by using macroturbulent laws is not always compatible with the $V\\,\\sin\\,i_\\star$, which is directly obtained from the Rossiter-McLaughlin measurement. ",
        "introduction": "The WASP consortium (Wide Angle Search for Planets) \\citep{Pollacco:2006fj} has been operating from La Palma, Spain, and Sutherland, South Africa. Its main goal is to find transiting extrasolar planets. With more than 80 planets discovered, this is the most successful ground-based survey for finding short-period giant planets. Amongst the many planet candidates that WASP has produced are many `false positives', which here we regard as objects of interest, that have been shown by radial-velocity follow-up to be M dwarfs that eclipse F, G or K-dwarf companions. They are of a few Jovian radii in size and thus mimic a planetary transit signal very well. Because of the mass and low brightness of the secondaries,  they remain invisible, making them convenient objects for follow-up and study using the same photometry and radial-velocity techniques that are routinely used for exoplanets. Two A+M binaries have already been presented in \\citet{Bentley:2009lr} and similar objects have been found by the OGLE survey \\citep{Udalski:2007lr,Pont:2006lr} and by the HAT network \\citep{Fernandez:2009fk}.  We have made a substantial effort to characterise these low-mass eclipsing binaries (the EBLM Project) in order to discover transiting brown dwarfs (such as WASP-30b \\citep{Anderson:2011fk}) and also to complete the largely empty mass--radius diagram for stars with masses $< 0.4$ M$_\\odot$. These objects explore the mass distribution separating stars from planets, or serve as extended samples to the exoplanets, especially with regards to their orbital parameters, long term variability and spin-orbit angles. Our results will be published in a series of papers, of which this is the first.\\\\   A primary goal of the EBLM Project is to address the M-dwarf radius  problem whereby current stellar evolution models underestimate the radii of M~dwarfs by at least 5\\% and overestimate their temperatures by a few hundred degrees (e.g. \\citep{Morales:2010fk, Morales:2009fj,Lopez-Morales:2007kx} and references therein). Thus we aim to substantially increase the number of M dwarfs with accurate masses, radii, and metallicities using a large sample of newly discovered eclipsing binaries comprising F, G, K primaries with M dwarf secondaries. The masses and radii results are inferred using F, G, K atmospheric and evolution models. Although model-dependent, the analysis of bright F, G, K + M dwarf eclipsing binaries promises large numbers of masses and radii of low-mass stars over the entire range of M dwarfs down to the hydrogen-burning limit. They will have accurate metallicity determination, and cover a wide range of activity levels. A combined analysis of the radial-velocity curve and light curve permits to deduce the masses and radii, while an accurate system metallicity can be determined from the F, G, K primary star. Furthermore, activity can be determined indirectly through knowledge of the rotation-activity relation \\citep{Morales:2008qy} combined with $V\\,\\sin\\,i_\\star$ from measurements or by deduction when the systems are tidally synchronised.\\\\  %One of the most compelling problems to arise in the last few years is the reported discrepancy between the observed and theoretically predicted radii and temperatures of %M~dwarf stars.       This problem manifests itself in many different objects  including both short and long period eclipsing binaries, interferometric radius measurements of single low mass stars, and in the colours of active field M dwarfs. %One goal of the EBLM Project, is to address this M~dwarf radius problem by substantially increasing the number of M~dwarfs with accurate masses, radii, and metallicities using a large sample of  newly discovered eclipsing binaries comprising F, G, K primaries with M~dwarf secondaries.  The analysis of bright F, G, K + M dwarf eclipsing binaries promises large numbers of masses and radii of low-mass stars over the entire range of M~dwarfs down to the hydrogen burning limit and with accurate metallicity determinations and a range of activity levels.   A combined analysis of the radial velocity curve and light curve permits direct measurement of the masses and radii, and an accurate metallicity can be determined from the relatively simple spectrum of the F, G, K primary star. Furthermore, activity can be determined indirectly through knowledge of the rotation-activity relation combined with $V\\,\\sin\\,i_\\star$ measurements or when the systems are tidally synchronized.  \\\\  \\citet{Holt:1893fk}, in proposing a method to determine the rotation of stars prior to any knowledge about line broadening, predicted that when one star of a binary eclipsed the other it would first cover the advancing blue-shifted hemisphere and then the receding red-shifted part. This motion would create a colour anomaly perceived as a progressive red-shift of the primary's spectrum followed by a blue-shift, thus appearing as a symmetric radial-velocity anomaly on top of the main Doppler orbital motion of the eclipsed star's lines. This effect was first observed by \\citet{Rossiter:1924qy} and \\citet{McLaughlin:1924uq}, though with some evidence of its presence noted earlier by \\citet{Schlesinger:1910fj} (p134). Holt's  idea was correct but only under the assumption that both stars orbit in each other's equatorial plane. In the case of a non-coplanar orbital motion the radial velocity effect is asymmetric (see e.g. \\citet{Gimenez:2006kx} or \\citet{Gaudi:2007vn})  Recent observations  of this effect in transiting extrasolar planets (e.g. \\citet{Queloz:2000rt,Winn:2005ys,Hebrard:2008mz,Winn:2009lr,Triaud:2010fr,Moutou:2011cr,Brown:2012lr} and references therein) have shown that the so-called \\textit{hot Jupiters}, gas giant planets on orbits $<$ 5 days, have orbital spins on a large variety of angles with respect to the stellar spin axis,  the most dramatic cases being on retrograde orbits. While it was previously thought that hot Jupiters had migrated from their formation location to their current orbits via an exchange of angular momentum with the protoplanetary disc, they are now thought to have been dynamically deflected onto highly eccentric orbits that then circularised via tidal friction. There are various ways in achieving this, such as planet--planet scattering \\citep{Rasio:1996ly,Nagasawa:2008gf,Wu:2011ul} and Kozai resonances \\citep{Kozai:1962qf,Wu:2007ve,Fabrycky:2007pd,Naoz:2011bh}. These could be triggered by environmental effects in their original birth clusters such as fly-bys \\citep{Malmberg:2007lq,Malmberg:2011dq}, by an additional, late, inhomogenous mass collapses in young systems \\citep{Thies:2011yq}, or during the planet formation process itself \\citep{Matsumura:2010ve,Matsumura:2010ul}.  Several patterns have emerged in the planetary spin--orbit angle data, including: a lack of aligned systems whose host stars have $T_\\mathrm{eff} > 6250$ K \\citep{Winn:2010rr}; a lack of inclined systems older than 2.5--3 Gyr \\citep{Triaud:2011fk}; and a lack of retrograde system for secondaries $> 5$ M$_\\mathrm{Jup}$ \\citep{Hebrard:2011fk,Moutou:2011cr}. To help confirm this latter trend, one could measure the Rossiter--McLaughlin effect in several heavy planets, but those are rare. It is thus easier to extend the mass range to low-mass stars, hoping to further our understanding of the planetary spin--orbit angle distribution. \\\\  The fact that hot Jupiters can be on inclined orbits raises the question about the inclinations of close binary stars. As proposed by \\citet{Mazeh:1979eu},  close binaries, especially those with large mass differences, might form via the same dynamical processes that have been proposed for hot Jupiters, i.e., gravitational scattering followed by tidal friction. In fact, \\citet{Fabrycky:2007pd} primarily address the formation of close binaries; the possible application to exoplanets comes later. That paper was motivated by observational results, notably presented by \\citet{Tokovinin:2006la}, showing that at least 96\\,\\% of close binaries are accompanied by a tertiary component, supporting the appeal to the Kozai mechanism as described in \\citet{Mazeh:1979eu}. It has been argued that objects as small as 5 M$_\\mathrm{jup}$ could be formed as stars do \\citep{Caballero:2007mz}, while objects as massive as 20 or 30~M$_\\mathrm{jup}$ could be created by core collapse, in the fashion expected for planets \\citep{Mordasini:2009gf}. Rossiter--McLaughlin observations bridging the mass gap between planets and stars could eventually help in separating or confirming both proposed scenarii.\\\\  Even though attempts have been made to model the Rossiter--McLaughlin effect (e.g. \\citet{Kopal:1942ai,Hosokawa:1953kl}) no systematic, quantified and unbiased survey of the projected spin--orbit angle $\\beta$ in binary  star systems can be found in the literature. Only isolated observations of nearly aligned systems have been reported. \\citet{Kopal:1942qe} mentions a possibly asymmetric Rossiter--McLaughlin effect (or rotation effect as it was then known) leading to an estimated misaligned angle of $15^{\\circ}$ observed in 1923 in the Algol system, but that was presented as aligned by \\citet{McLaughlin:1924uq}. \\citet{Struve:1950tg} (p 125) writes that the rotation effect had been observed in a 100 systems without citing anyone. \\citet{Slettebak:1985dp} is a good source of citations about this epoch.  \\citet{Worek:1996hc} and \\citet{Hube:1982sp} are two examples of more recent observations of the Rossiter--McLaughlin effect. The rotation effect was also used for cataclysmic variables to determine if the accreting material comes from a disc in a  plane similar to the binary's orbital plane \\citep{Young:1980ij}.  It has to be noted that, early on, the Rossiter--McLaughlin effect was used as a tool to measure the true rotation of a star, hence creating a bias against reporting misaligned systems. Furthermore the precision and accuracy of instrumentation, data extraction and analysing technique of that time prevented the observation of the Rossiter--McLaughlin effect for slowly rotating stars, further biasing detections of the effect towards synchronously rotating binaries, which could have tidally evolved to become aligned \\citep{Hut:1981kx}.  In addition, the capacity to accurately model blended absorption lines of double-lined binaries (SB2) during transit has only been developed recently. Thus, most people that studied binaries chose not to observe during eclipses. Modelling eclipsing SB2 has recently been developed in \\citet{Albrecht:2007th}, and used by \\citet{Albrecht:2009fy} for the case of DI Herculis, explaining its previously abnormal apsidal motion: both stars orbit above each other's poles. These measurements are being followed by a systematic and quantified survey of spin--orbit measurements for SB2s of hot stars with similar masses (the BANANA project, \\citet{Albrecht:2011bs}). Another contemporary result is presented in an asteroseismologic paper by \\citet{Desmet:2010fv}.  We circumvent the  blended-line problem altogether by only observing the WASP candidates that turned out to be single-line binaries (SB1) while searching for extrasolar planets. Low-mass M dwarfs and brown dwarfs have a size similar to gas giants and appear to a first approximation like a planet: a dark spot moving over the disc of their primary. Thus, the low-mass eclipsing binaries found by transiting planet surveys  provide a  good sample to extend the work carried out on planets and provide a largely unbiased, quantified survey of spin--orbit angles for F, G or K~+~M binaries, complementary to the BANANA project.  The differences between our primaries will also allow us to probe the way tides propagate in convective or radiative stars \\citep{Zahn:1977yq}. In stellar parameters and data treatment, our systems resemble the aligned pair Kepler-16 A\\,\\&\\,B \\citep{Doyle:2011vn,Winn:2011wd}, but with shorter periods.  In this we first present our observations of WASP-30 and J1219--39 (1SWASPJ121921.03--395125.6, TYC 7760-484-1), then describe our models and how they were adjusted to fit the data, and how the error bars were estimated. We will then move to a discussion of the results. %Another method to determine the angle of close binaries consists in doing Doppler imaging as explained in \\citet{Strassmeier:2000p7878} ",
        "conclusions": "We announce the discovery of a new low-mass star whose mass and radius have been precisely measured and found to be at the junction between the stellar and substellar regimes. In addition, using observations with the CORALIE spectrograph, on the 1.2\\,m Euler Telescope, and HARPS, on the ESO 3.6\\,m, we have demonstrated the detection of the Rossiter--McLaughlin effect on two objects more massive than planets. These measurements are amongst the first to be realised on such objects. They will help study the dynamical events that could have led to the formation of binary systems where both components have a large mass difference, and may also provide a useful comparison sample to the spin--orbit angle distribution of hot Jupiters, as well as helping theoretical developments in the treatment of  tides, the main mechanism behind synchronisation, circularisation and realignment.   \\citet{Hebrard:2011fk} and \\citet{Moutou:2011cr} note that while hot Jupiters are usually found with a large variety of orbital angles, objects above 5 $M_\\mathrm{jup}$ are not found on retrograde orbits. We extend the distribution of spin--orbit angle versus mass to beyond the planetary range. Both of our objects appear well aligned with their primary, further confirming that trend. Being around stars colder than 6250\\,K, they also reinforce the pattern shown by \\citet{Winn:2010rr} between orbital inclination and the primary's effective temperature. The age of WASP-30A and the alignment of WASP-30b also helps strengthen the pattern claimed in \\citet{Triaud:2011fk} that systems older than 2.5--3 Gyr are predominantly aligned. It is interesting to note that while J1219--39b is on a slightly eccentric orbit, its orbital spin is aligned with its primary's rotational spin. This gives observational evidence that orbital realignment may be faster than orbital circularisation for these objects, in opposition with planets. Final tidal equilibrium has not been reached for either system as WASP-30A is not synchronised and J1219--39b is not circularised (nor synchronised) \\citep{Hut:1980bh}.\\\\  \\begin{figure} \\centering \\includegraphics[width=9cm]{massrad2.pdf} \\caption{Mass--radius diagram for heavy planets, brown dwarfs and low-mass stars. The radius axis corresponds to the size range of Jupiter-mass planets discovered so far. Inverted blue triangles show eclipsing/transiting SB1s, upright red triangles denote interferometric measurements. The two (M$_2$, R$_2$) posterior probability density distributions for WASP-30b and J1219--39b are drawn in grey with their $1-\\sigma$ confidence regions in white. Models by \\citet{Baraffe:2003gf,Baraffe:1998ly} are also displayed with ages 5\\,Gyr (black), 1\\,Gyr (dark grey), 500\\,Myr (light grey) and 10\\,Gyr (dotted). Models are for [M/H] $= 0 $. Observational data were taken from \\citet{Lane:2001kx,Segransan:2003fr,Pont:2005qy,Pont:2005fk,Pont:2006lr,Beatty:2007uq,Deleuil:2008lr,Demory:2009ys,Bouchy:2011lr,Johnson:2011fk,Ofir:2012lr,Siverd:2012fk}.} \\label{fig:mr} \\end{figure}   While not being the primary objective of this paper, we present a method to analyse, in a global manner, eclipse photometry, the radial-velocity reflex motion, and the Rossiter--McLaughlin effect, for objects more massive than planets, in order to obtain precise estimates of the mass, radius and orbital parameters of SB1s. The precision of a few percent that we obtain comes from our use of $\\rho_\\star$, the mean stellar density, instead of the more traditional $\\log\\,g$ when interpolating inside the stellar evolution tracks. This interpolation gives us precise values for the primary's stellar parameters which are used to estimate the secondary's parameters. We can check our method by deducing $\\log\\,g_\\star$ and comparing them to their spectral counterpart. The values are in very good agreement and fall within the errors of the spectral method (see table \\ref{sec:results}). This method also allows us to estimate ages from reading the stellar tracks, something important in the case of close binaries where gyrochronology cannot be trusted owing to the tidal evolution of the system.  Our low-mass eclipsing objects have very similar surface gravities, but, located on opposite sides of the brown-dwarf limit, have sizes dominated by different physics \\citep{Baraffe:2003gf}. Plotting the posterior probability distribution for the mass and radius on theoretical mass--radius relations \\citep{Baraffe:2003gf,Baraffe:1998ly} in Figure \\ref{fig:mr}, we observe that WASP-30b is between the 0.5 and 1 Gyr tracks, suggesting the object is fairly young (and thus luminous, which should cause a measurable secondary eclipse). This is at odds with the age of the primary, which we have found to be older than 2 Gyr with $99\\%$ confidence, with a best age of $3.4\\pm0.4$ Gyr (figure \\ref{fig:w30age}).  We could explain this radius anomaly if the object has been inflated in the same manner that has been observed for hot Jupiters thanks to the high irradiation received from its primary \\citep{Demory:2011lr}. The exact physical causes are still being debated. It may also be that energy has been stored inside the object if it circularised from a previously highly eccentric orbit \\citep{Mazeh:1979eu,Fabrycky:2007pd}.  WASP-30b's mass is interesting in that it is close to the minimum of the tracks presented in \\citet{Baraffe:2003gf} and displayed in figure \\ref{fig:mr}. The mass--radius posterior distribution of J1219--39b is compatible with the 10-Gyr theoretical line. Better photometry would help to reduce the confidence region. This analysis will be the subject of subsequent papers (Hebb et al. in prep). One could object that our analysis does not take into account the fact that both WASP-30b and J1219--39b are self luminous. Even in the case that they had the same effective temperature as their primaries, the overall contamination cannot be larger than their relative sizes $\\sim 1\\,\\%$, lower than our current precision.  We would like to attract the attention on the fact that both those objects have sizes entirely compatible with those of hot Jupiters. While hot Jupiters are often inflated, Jupiter-mass planets at longer periods no longer are \\citep{Demory:2011lr}. It would then be expected that  many of the planet candidates published by the space mission \\textit{Kepler}  with inflated radii ($> 1.2\\,R_\\mathrm{Jup}$) and periods longer than 10--15 days could be objects similar to WASP-30b and J1219--39b. %In principal, those objects could be identified by measuring their secondary occultation, but, as J1219-39b suggest, after a 5 day orbit, low mass binaries start being on eccentric orbits meaning the chance of having both the primary and secondary eclipses is lowered. It may even be that what \\textit{Kepler} has,  in some cases, detected  the secondary eclipse instead of the primary. While not being planets, they are of great interest for their masses, radii and orbital parameters.\\\\   Finally, the Rossiter--McLaughlin effect has recently been used almost exclusively to measure planetary orbital planes. Observing it for binary stars extends that work by bridging the gap in mass ratio between planetary and stellar systems. Comparison between low-mass binaries in our case with higher-mass binaries as in the BANANA survey \\citep{Albrecht:2011bs} will permit us to test different regimes of binary formation and tidal interactions.  Yet more information still lies in the study of this RV anomaly. Its use in the beginning of the $20^\\mathrm{th}$ century was primarily to measure the rotation of stars. As seen in this paper, there is a discrepancy between the \\vsini values obtained by using calibration of the macroturbulence and the directly measured $V\\,\\sin\\,i_\\star$ from the Rossiter--McLaughlin effect. This was also pointed out in \\citet{Triaud:2011vn} in the case of WASP-23, and by \\citet{Brown:2012lr} in the case of WASP-16. Both those systems, like J1291--39, contain K dwarfs, and both had their spectroscopic $v\\,\\sin\\,i_\\star$ overestimated compared to the value obtained via the Rossiter--McLaughlin effect. Since we compute the spin--orbit angle $\\beta$ and the $V\\,\\sin\\,i_\\star$, we have strong constraints on the coplanarity of the system and rotation velocity of the primary; thanks to the transit/eclipse geometry, we obtain accurate and precise masses and radii for both the primary and the secondary. Combining all this information and collecting many measurements, we will be able to test which of the macroturbulence laws one should use. Should none apply, inserting the observed $V\\,\\sin\\,i_\\star$ values as input parameters in spectral line analyses we will have the capacity to measure macroturbulence directly. Observing the Rossiter--McLaughlin effect is thus not just about glimpsing into the past dynamical history of systems, but can also become an important tool for  understanding  stellar physics better.  \\paragraph{\\textbf{Nota Bene}} We used the UTC time standard and Barycentric Julian Dates in our analysis. Our results are based on the equatorial solar and jovian radii and masses taken from Allen's Astrophysical Quantities."
    },
    "1208/1208.1409_arXiv.txt": {
        "abstract": "We present a  study of the kinematics of the isolated spiral galaxy NGC 864, using H$ \\alpha $ Fabry-Perot data obtained with the Galaxy H$ \\alpha $ Fabry-Perot System (GH$ \\alpha $FaS) instrument at the William Herschel Telescope in La Palma, complemented with images at 3.6 $ \\mu  $m, in the \\textit{R} band and in H$ \\alpha $ filter, and integral-field spectroscopic data. The resulting data cubes and velocity maps allow the study of the kinematics of the galaxy, including in-depth investigations of the rotation curve, velocity moment maps, velocity residual maps, gradient maps and position-velocity diagrams. We find asymmetries in the velocity field in the bar zone, caused by non-circular motions, probably in response to the potential of the bar. We also find a flat-profile bar, in agreement with the strong bar, with the grand design spiral pattern, and with the gap between the ends of the bar and the start of the spiral arms. We quantify the rate of massive star formation, which is concentrated in the two spiral arms. ",
        "introduction": "Galaxies are the basic building blocks of the Universe and their formation and evolution are of great interest in current astrophysical research. Their dynamics and morphology are the result of both externally driven (e.g. galaxy mergers) and secular evolution (e.g. bar or spiral pattern driven). Disentangling the different early evolutionary tracks is complicated by redshift, distance and dust absorption. Alternatively one can explore the fossil record of evolution through detailed observations of the end product, nearby galaxies.  One of the most important drivers of this internal secular evolution is the flow of gas into the central regions. This is the result of angular momentum loss in shocks induced by non-axisymmetric potentials, as generated by a bar, interactions, minor mergers or even by minor deviations from axisymmetry such as ovals, lenses, lopsidedness or spiral arms (e.g. \\citealt{Schwarz1984}; \\citealt{Shlosman1989,Shlosman1990}; \\citealt{Knapen1995}; \\citealt{Rix1995};   \\citealt{Kormendy2004}; \\citealt{Comeron2010}). Around two-thirds of local galaxies exhibit at least one bar (\\citealt{RC3}; \\citealt{Sellwood1993}; \\citealt{Moles1995}; \\citealt{Ho1997}; \\citealt{Mulchaey1997}, \\citealt{Hunt1999}; \\citealt{Knapen2000}; \\citealt{Eskridge2000}; \\citealt{Laine2002}; \\citealt{Laurikainen2004a}; \\citealt{Menendez-Delmestre2007}; \\citealt{Marinova2007}; \\citealt{Sheth2008}; \\citealt{Laurikainen2009}), so potential drivers of secular evolution are common. Low-mass discs acquired their bars later than more massive ones (\\citealt{Sheth2008}), implying that the bar-driven secular evolution is expected to be different depending on each galaxy's mass.  To study the influence of the past evolution of a galaxy on the observed morphology, kinematic information is essential. The study of velocity fields in spirals has mostly been done using the 21 -cm HI line, primarily because this emission can be traced far out, often three or four times beyond the visible disc. \\citet{Rots1975}, \\citet{vanderHulst1979}, \\citet{Bosma1981},  \\citet{Gottesman1982} and others demonstrated the power of this technique in deriving the total mass distribution of disc galaxies. Since the beam sizes associated with 21 cm observations were typically at least as large as the disc scale length of the galaxy, little small-scale structure in the velocity field was detected in this manner. The HI Nearby Galaxy Survey (THINGS; \\citealt{Walter2008}) presents some of the highest spatial resolution ($\\sim6''$) 21-cm HI observations of nearby galaxies to date, using the Very Large Array (VLA) of the National Radio Astronomy Observatory (NRAO).  The best candidate technique to yield two-dimensional kinematical maps across a whole galaxy with good arcsecond spatial resolution is H$\\alpha$. Due to the cosmic abundance of hydrogen, H$\\alpha$ is often the brightest emission line in the visible wavelength range. In spiral galaxies, this line traces primarily the ionized gas in HII regions around young massive stars. Traditional long-slit spectra with optical emission lines are valuable for deducing the rotation curve of a galaxy (e.g., \\citealt{Rubin1980}, \\citealt{Rubin1982a}, \\citealt{Rubin1985}, \\citealt{Amram1992}, \\citealt{Amram1994}, or \\citealt{Persic1995}), and the rotation curves can be extended to very large radii using H$ \\alpha $ (e.g. \\citealt{Christlein2008}). H$\\alpha$ observations using Fabry-Perot (FP) instruments have been used for some 40 years now (\\citealt{Tully1974}; \\citealt{Deharveng1975}, \\citealt{Dubout1976} or \\citealt{deVaucouleurs1980}), and have since been employed to create high signal-to-noise ratio (S/N) rotation curves for many spiral galaxies (e.g., \\citealt{Marcelin1985}; \\citealt{Bonnarel1988}; \\citealt{Pence1990}; \\citealt{Corradi1991}; \\citealt{Cecil1992}; \\citealt{Sicotte1996}; \\citealt{Ryder1998}; \\citealt{Jimenez-Vicente1999}; \\citealt{Garrido2002}, \\citealt{Epinat2008}, and many more).  Using the Infrared Array Camera (IRAC; \\citealt{Fazio2004}) operating at 3.6 and 4.5 $ \\mu $m, the \\textit{Spitzer} Survey of Stellar Structure in Galaxies (S$^4$G; \\citealt{Sheth2010}) targets over 2300 galaxies. The sample is volume-, magnitude- and size-limited (\\textit{d}$<$40 Mpc, \\textit{$m_b$}$<$15.5, \\textit{$D_{25}$}$>$1 arcmin). The cornerstone of the survey is the quantitative analysis of photometric parameters, enabling a variety of studies on secular evolution, outer disc and halo formation, galaxy morphology etc. The new images, much deeper than traditional ground-based near-IR observations, will allow a comprehensive and definitive study of galaxy structure not only as a function of stellar mass, but also as a function of environment, vital to test cosmological simulations predicting the mass properties of present day galaxies.  We have designed an observing programme to obtain H$\\alpha$ FP kinematic data sets of over 40 spiral galaxies that are included in the S$ ^{4} $G sample. The observation of NGC 864 is part of this wider programme using the instrument Galaxy H$ \\alpha $ Fabry-Perot System (GH$ \\alpha $FaS) at the William Herschel Telescope (WHT) in La Palma. NGC 864 is used to illustrate the data and methods, and discuss the kinds of results the main survey will give. ",
        "conclusions": "We have presented a kinematic study of the spiral galaxy NGC 864. Data sets obtained with three different WHT instruments have been studied, resulting in rotation curves and intensity, velocity, residual and gradient maps. From this, we reach the following conclusions:  \\begin{enumerate} \\item A technique for flux calibration of GH$ \\alpha $FaS data has been developed. This will be used for other S$ ^{4} $G galaxies observed with GH$ \\alpha $FaS. \\item Our bulge/disc/bar 2D decomposition of the S$^4$G 3.6$\\mu m$ image of NGC 864 using {\\sc budda} yields a flat-profile bar, which is typical of strong bars (\\citealt{Athanassoula2002}) and of two-armed grand design spirals (\\citealt{Elmegreen1985}). Flat-profile bars are linked to a higher degree of exchange of angular momentum with the halo than exponential bars (\\citealt{Athanassoula2003}). Such bars can result from sharply rising rotation curves, as observed in NGC 864, and are most likely connected to the observed spatial offset between the ends of the bars and the onset of the spiral arms. \\item The rotation curves obtained from the GH$ \\alpha $FaS velocity maps are similar to those obtained by \\citet{Epinat2008}, both using H$ \\alpha $ observations. However, the curves do not cover all the disc and the intrinsic patchiness linked to the line emission prejudices the analysis and the derived results. Due to the high angular and spectral resolution, our new rotation curve shows considerable detail in the central 2 kpc radius. \\item We have found non-circular motions along the bar in the residual maps, in the PV diagram along the kinematic minor axis, and in the velocity profile along the kinematic minor axis. These velocity patterns are typical of a barred galaxy behaviour, with streaming motions along the bar.  \\end{enumerate}  This paper is the first in a series presenting the results of the programme to observe a sample of some 40 spiral galaxies of all morphological types from the S$ ^{4} $G sample with the GH$ \\alpha $FaS FP instrument."
    },
    "1208/1208.4689_arXiv.txt": {
        "abstract": "The time-scale over which and modality by which young stellar objects (YSOs) disperse their circumstellar discs dramatically influences the eventual formation and evolution of planetary systems. By means of extensive radiative transfer (RT) modelling, we have developed a new set of diagnostic diagrams in the infrared colour-colour plane (K\\,-\\,[24] vs. K\\,-\\,[8]), to aid with the classiffication of the evolutionary stage of YSOs from photometric observations. Our diagrams allow the differentiation of sources with un-evolved (primordial) discs from those evolving according to different clearing scenarios (e.\\,g.\\,homologous depletion vs. inside-out dispersal), as well as from sources that have already lost their disc. Classification of over 1500 sources in 15 nearby star-forming regions reveals that approximately 39\\,\\% of the sources lie in the primordial disc region, whereas between 31\\,\\% and 32\\,\\% disperse from the inside-out and up to 22\\,\\% of the sources have already lost their disc. Less than 2\\,\\% of the objects in our sample lie in the homogeneous draining regime. Time-scales for the transition phase are estimated to be typically a few $10^{5}$ years independent of stellar mass. Therefore, regardless of spectral type, we conclude that currently available infrared photometric surveys point to fast (of order 10\\,\\% of the global disc lifetime) inside-out clearing as the preferred mode of disc dispersal. ",
        "introduction": "The lifetime and modality for the dispersal of protoplanetary discs around newly formed low mass stars (approximately solar mass or lower) is a key parameter that influences the formation and evolution of eventual planetary systems. The classical picture that emerged from the last decade of photometric observations, mainly carried out with the Spitzer Space telescope is that of disc evolution being described by two different timescales. The first timescale could be defined as a \u2019global timescale\u2019, i.\\,e.\\,the total time it takes from a star to go from disc-bearing to disc-less, and a \u2019dispersal timescale\u2019, i.\\,e.\\,the time it takes for a disc to disappear once dispersal has set in. Global timescales, which can be inferred from the study of disc frequencies (e.\\,g.\\,\\citet{Haisch2001}), are of order a few million years (e.\\,g.\\,\\citet{Hernandez2007b, Mamajek2009}). Dispersal time-scales, as determined from the study of infrared colours of young stars, appear to be much shorter, indicating that the dispersal mechanism must be fast (e.\\,g.\\,\\citet{Kenyon1995}; \\citet{Luhman2010}; \\citet{Ercolano2010}). Such observed two-timescale behaviour has favoured the development of disc dispersal models that involve a rapid disc clearing phase, contrary to the predictions of simple viscous draining, and in agreement with photoevaporation \\citep{Clarke2001, Alexander2006a, Alexander2006b, Ercolano2008, Ercolano2009, Gorti2009, Owen2010, Owen2011b, Owen2011a, Owen2012a} or possibly planet formation \\citep{Armitage1999}. The interpretation of infrared colours in relation to the evolutionary state of a disc is, however, far from being trivial. This is particularly true with regards to the classification of transition discs, the latter being intended as objects caught in the act of disc dispersal. The evolution of the dust component in a disc is mirrored by the evolution of colours in the infrared plane. By means of radiative transfer modelling, \\citeauthor*{Ercolano2010} (2011, henceforth \\citetalias{Ercolano2010}) identified the regions in the K\\,-\\,[8] vs. K\\,-\\,[24] plane where primordial discs, discs with inner-holes (i.\\,e.\\,presumably being dispersed from the inside-out) and discs which lose mass homogeneously at all radii, are expected to be found. Their study, which was limited to M-stars, showed that in the case of the cluster IC348, most discs disperse from the inside-out and undergo the transition on a short time-scale, as predicted by standard photoionisation models. These conclusions are in contrast with the conclusions of \\citet{Currie2009b}, who claimed instead a large number of \u2019homogeneously depleting\u2019 discs, for the same cluster. Such discrepancies highlight the need for detailed modelling in the interpretation of IR colours of discs. The study of \\citetalias{Ercolano2010} was restricted to M-stars in only one cluster, which prevented the authors from being able to make a more general statement with regards to disc dispersal. Here we present results from a forthcoming paper \\citep{Koepferl2012}, which significantly improves on the work of \\citetalias{Ercolano2010} by performing further RT calculation to evaluate evolutionary tracks in the IR colour plane for stars of different spectral types. We then apply our results to the photometric data of 15 nearby star-forming regions, that we collected from the literature, in order to address the question of what is the preferred mode of disc dispersal. ",
        "conclusions": "We have calculated the SEDs of protoplanetary discs of different spectral types, geometries, settling and inclination. We then considered the evolution of the infrared colours (K\\,-\\,[8] vs. K\\,-\\,[24]) of the model disks as they disperse according to different scenarios (homologous depletion, inside-out and outside-in clearing). Based on our models we propose a new diagnostic infrared colour-colour diagram to classify the evolutionary stage of YSOs. We have applied our infrared colour-colour diagnostic diagram to classify YSOs in 15 nearby star-forming regions and study the evolution of their disc populations. We estimate time-scales for transition phase of typically a few $10^{5}$ years independent of stellar mass. We conclude that, regardless of spectral type, current observations point to fast inside-out clearing as the preferred mode of disc dispersal."
    },
    "1208/1208.1779_arXiv.txt": {
        "abstract": "We study the short--term topological changes of equatorial and polar coronal hole (CH) boundaries, such as a variation of their area and disintegration, associated to reconnection with nearby (within 15$^\\circ$ distance) quiescent prominence magnetic fields leading to eruptions and subsequent Coronal Mass Ejections (CMEs). The examples presented here correspond to the recent solar minimum years 2008 and 2009. We consider a temporal window of one day between the CH topological changes and the start and end times of prominence eruptions and onset of CMEs. To establish this association we took into account observational conditions related to the instability of prominence/filaments, the occurrence of a CME, as well as the subsequent evolution after the CME. We found an association between short--term local topological changes in CH boundaries and the formation/disappearance of bright points near them, as well as, between short--term topological changes within the whole CH and eruptions of nearby quiescent prominences followed by the appearance of one or more CMEs. ",
        "introduction": "Coronal holes (CHs) are areas which are seen dark in X--rays and extreme ultraviolet \\citep{waldmeier1,bohlin1,wang1}. Their plasma temperature and density is much lower than that of the ambient corona; so, they are low emission zones \\citep{harvey2,raju1}. One of the important characteristics of CHs is their magnetic field, which is characterized by being mainly unipolar \\citep{wang1,harvey2,raju1} with open magnetic field lines \\citep{bohlin1,wang1,harvey2}; this means that its direction is primarily radial and the tangential component can be considered zero for distances greater than 2.5Rs \\citep{wang3}, with Rs being the solar radius.  Closed magnetic field lines can keep the plasma within coronal loops; however, since the field is open in CHs, the plasma can escape into the interplanetary medium making up the solar wind.   It has been shown that the emergence of a small bipolar region from below the solar surface and its interaction with the preexisting open field in the coronal hole is a prime candidate to trigger reconnection and the consequent launching of jets along field lines in the corona \\citep{moreno}.      Magnetic reconnection, which may occur between open and closed magnetic field lines or between open lines \\citep{kahler2} is of great importance in the study of CHs, especially near their boundaries.  The process of continuous reconnection at CH boundaries is known as ``interchange reconnection'' \\citep{wang4,wang1,fisk2,raju1,edmondson}.  \\citet{wang2} have proposed two kinds of interchange reconnection, the first occurs when two open field regions of the same polarity are separated by photospheric flux of the opposite polarity and the reconnection takes place between the open flux domains and the underlying pair of loop systems. The second involves stepwise displacements within a region of single magnetic polarity and occurs when an open field line exchanges footpoints with a closed field line rooted next to it and the reconnection happens at the apex of the closed loop. Hence, coronal magnetic reconnection near CH boundaries can be responsible for their shape, magnetic topology and evolution.  Moreover, the appearance/disappearance of bright points (BPs) associated with short term evolution of CH boundaries \\citep{nolte1,davis1,kahler1,madjarska1}, as well as CH topological changes associated with CMEs \\citep{gonzalez1} or with filament eruptions and subsequent CMEs \\citep{bravo1,gopalswamy2,jiang1,taliashvili2} reported in previous studies can be explained by the associated magnetic field configuration.    Regarding the origin of CMEs, statistical studies show that 70\\% of prominence eruptions are associated with CMEs \\citep{munro,pojoga}, moreover CHs may also be considered as signatures of CMEs \\citep{bravo2,thompson}. In addition, filament eruptions are associated with the formation of small adjacent CHs and/or topological fluctuations of CHs \\citep{harvey1,taliashvili1,taliashvili2}; and/or increased CH area \\citep{bravo2,taliashvili1,taliashvili2}; and/or decrease and/or its complete disappearance \\citep{taliashvili1,taliashvili2}. Using H$\\alpha$ movies and spectroheliograms taken at Observatoire de Paris-Meudon, \\citet{taliashvili2} reviewed 42 quiescent solar filaments/ prominences eruptions during two minimum periods of solar activity (1985--1986 and 1994). These authors found that the majority of prominence/filament eruptions ($\\sim$91\\%) are associated with the presence of adjacent CHs and subsequent CMEs. Magnetic reconnection is probably the mechanism responsible for the interaction between CHs and prominences located close to their boundaries, which may result in the onset of CMEs. However, the characteristics of this process is not yet clear. Recent studies show that nearly 15$^\\circ$ distance between the prominences and CH boundaries could be considered as a possible critical distance for their interaction \\citep{taliashvili2}.  The Sun is the primary driver of space weather; the most severe solar activity occurs around the cycle maximum, but its influence on the magnetosphere and ionosphere continues through the solar minimum \\citep{cole}. Major events occur due to CMEs and high-speed solar wind streams associated with CHs \\citep{feldman,echer,cole,schwenn}, as the solar wind speed originating in CHs is of about 750--800 km/s \\citep{wang1,raju2}. It is important to note that both the CHs as well as their offspring, the high-speed solar wind streams, are representative of an inactive or quiet Sun \\citep{feldman2,bame,schwenn}.  Recent studies of the slow decline of solar Cycle 23 and slow rise of Cycle 24 show that the low solar activity lasted from about 2006 to the end of 2009, with 2008 and 2009 being particularly quiet years \\citep{toma}, which is reflected in several solar irradiance observations, such as solar UV and EUV irradiance, and radio flux at 10.7 cm wavelengths \\citep{tsurutani,solomon,hathaway}. Moreover, in 2007 and 2008, the effect of multiple, large, and long-lived CHs resulted in regular and recurrent solar wind streams and fast wind above 550 km/s, whereas in 2009, the disappearance of the low-latitude CHs shifted the sources of the solar wind to higher latitudes, mostly to the edges of the polar CHs with a drop of about 20\\% in the mean speed and the almost total disappearance of fast solar wind. Generally, the continuous presence of multiple, low--latitude CHs during 2007 and 2008 and the very low magnetic flux emergence, made the minimum between Cycles 23 and 24 different from the two previous minima \\citep{toma}.    In this work we study the possible association between CH topological changes and nearby prominence eruptions followed by CMEs as well as by sequences of CMEs. We analyze three events involving the evolution of the boundaries of CHs and of their surrounding regions (within $15^\\circ$ distance), taking a temporal window of 12 days. We identify their topological changes from one day before until one day after a CME occurrence. The events took place during the recent solar minimum years, 2008 and 2009. The short--term topological changes of studied CHs are accompanied by the disappearance of nearby quiescent filaments, as well as subsequent CMEs. ",
        "conclusions": "We study the detailed evolution of three different systems composed by coronal holes and nearby (within 15$^\\circ$ distance from their boundaries) quiescent filaments: CH1+F1, CH2+CH3+F2 and CH4+F3, that include polar CHs as well as equatorial ones. For each system we analyze the short--term CH topological changes during a period of 12 days associated with the disappearance of nearby filaments, especially before and/or after the disappearance. Moreover, we study the associated CMEs observed at similar position angles (in total 7 CMEs, excluding those with less than 10$^\\circ$ width) as the studied prominences and nearby coronal holes.  For all the studied events we observed the ejection of CMEs after the onset of the filament eruptions, within 36 hours of the starting time of prominence eruptions. We observe that during a prominence eruption at the northern hemisphere, a section of the CH (event 1) or the entire CH (event 2), located in opposite direction to this eruption, grows in extension; whereas the other section during the event 1, which is located towards the moving erupting prominence, decreases. Additionally, for event 2, the increase of the CH areas is more pronounced, probably due to the merging of both CHs. For both events, we observe associated bright points (two BPs for each events) before the starting time of prominence eruptions and their disappearance after the eruption. In addition, at least one of the BPs is associated to the visible reduction and fragmentation of the CH. In the case of event 1, the BP (located at the eastern section of CH1) lifetime coincides with the growth of this eastern section and with the period of time between the onset of CME1.1 and CME1.3, whereas the disappearance of the BP (located at the western section of CH1) coincides with the decrease of the western section and with CME1.2 starting time; after the onset of CME1.2, this western section disappears completely. For event 2, the appearance of one of the BPs coincides with the maximum separation between both CHs (CH2 and CH3); then, it disappears after the CME2 onset, whereas the lifetime of the other BP coincides with the time period between the disconnection and reconnection of the CHs (CH2+CH3) and it disappears before the CME2 onset.  Regarding event 3, located at the South Pole, we observe the appearance of a BP (BP1) before the starting time of prominence eruption followed by the appearance of other BP (BP2) and two CMEs (CME3.1, CME3.2), then BP1 intensifies, simultaneously CH4 fragments in two main sections around this BP; this is followed by CME3.3. After these 3 CMEs the BPs disappear. Each of the bright points observed near the CH (event 1 and event 2) boundaries is associated  with a small magnetic dipole that appears and disappears simultaneously with BP.  The most evident topological changes of the CHs are seen after the onset of the CMEs. For event 1 (North Pole CH), between the starting times of CME1.1 and CME1.3, the eastern section of CH1 (which lays in the same direction as CME1.1 and CME1.3) decreases its latitudinal width and increases its longitudinal width, whereas the small part of its western section disappears before the onset of the west limb CME1.2. For event 2, CH2 (North Pole CH) and CH3 (equatorial CH) appear separated after CME2 onset. Regarding event 3 (South Pole CH), CH4 disintegrates after CME3.2 onset, and after CME3.3 onset the separation between these segments reaches its maximum.  We have found that short--term topological changes in the entire CHs are associated with nearby quiescent prominence eruptions, particularly when a separation between the quiescent filaments and the CH boundary is less than or equal to 15$^\\circ$, followed by the ejection of one or more CMEs \\citep{taliashvili2}.  Similar results regarding the separation distance (of $\\sim$10$^\\circ$) between the filament and CH boundaries, associated with prominence eruption and subsequent CME were reported recently by \\citet{panasenco1}.   In addition, our observations indicate that BP appearance or disappearance is related with the starting processes of nearby filament eruptions, which can be caused by disturbances of the CH environment. These disturbances are most likely the result of magnetic reconnection between magnetic field lines associated with BPs and the surrounding of CH boundaries followed by the reorganization the magnetic field, which results in different BP lifetimes and CH boundary or area changes and reaches the foot points of nearby filaments, contributing in this way to their destabilization, eruption, and subsequent CME ejection. There are several proposed models and observational evidences related to magnetic reconnection at (or close to) the CH boundaries  and related with visible appearance/disappearance of BPs, filament eruptions, short-term topological changes in the entire CHs and subsequent formation of CMEs \\citep{kahler1,kahler2,bravo1,bravo2,madjarska1,wang2,fisk2}. Specially, magnetic bipoles emerging within a CH or small loops formed in the CH during interchange reconnection could increase along a CH lifetime and ultimately lead to the fragmentation and diffusion of the CH \\citep{krista1}. Interchange reconnection at CH boundaries is supported by several observational results \\citep{krista1}, e.g., the prevalence of very small loops inside CHs and larger loops outside CHs  \\citep{wiegelmann} and the boundary displacements observed due to the emergence and disappearance of bright small--scale loops in the form of BPs \\citep{madjarska2,subramanian}.  In addition, we observe that the direction of the erupted filaments near CHs and associated CMEs is almost non--radial. Both structures (regarding the three studied events) are moving toward the equator. In addition, the difference between the central position angles of both ranges between 15$^\\circ$ and 20$^\\circ$. These observations are consistent with the results reported by \\citet{gopalswamy1,cremades1,gopalswamy3,panasenco1};  they found that the CMEs generally move away from the open magnetic field regions and their the deflection is probably due to the fact that at lower coronal heights they are guided by the open field along which the fast solar wind flows.   In this study we are not considering long--term topological variation of coronal holes in regard to different stages of the solar cycle and the associated involvement of prominence eruptions and subsequent CMEs, which we plan to pursue in an upcoming project.    \\vspace{0.5cm}  {\\footnotesize{\\bf Acknowledgments.}We are grateful to the Hinode, STEREO, SOHO and Global High Resolution H--alpha Network for open access to their data sets. Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partner. It is operated by these agencies in co--operation with ESA and the NSC (Norway). LASCO and EIT are part of SOHO, SOHO is a project of international cooperation between ESA and NASA. The LASCO CME catalog is generated and maintained at the CDAW Data Center by NASA and The Catholic University of America in cooperation with the Naval Research Laboratory. The STEREO mission is supported by NASA, PPARC (UK), DRL (Germany), CNES (France), and USAF. The SECCHI data used here were produced by an international consortium of the Naval Research Laboratory (USA), Lockheed Martin Solar and Astrophysics Lab (USA), NASA Goddard Space Flight Center (USA), Rutherford Appleton Laboratory (UK), University of Birmingham (UK), Max--Planck--Institut for Solar System Research (Germany), Centre Spatiale de Li\\`ege (Belgium), Institut d'Optique Theorique et Appliqu\\'e (France), Institut d'Astrophysique Spatiale (France). The ``COR1 Preliminary Events List'' was generated by O. C. St. Cyr prior to September 2007, and is being maintained now by Hong Xie.  Wilcox Solar Observatory is currently supported by NASA and data used in this study was obtained via the web site, courtesy of J.T. Hoeksema. We are grateful to M. S\\'anchez and T. Roinishvili to improve English. This study was performed as a partial requirement for the PhD Degree of Sciences at the University of Costa Rica.  Special thanks are owed to anonymous referees for constructive comments that helped to improve the quality of the paper.}"
    },
    "1208/1208.3583_arXiv.txt": {
        "abstract": "The Kepler planet candidates are an interesting testbed for planet formation scenarios. We present results from $N$-body simulations of multi-planetary systems that resemble those observed by Kepler. We add both smooth (Type I/II) and stochastic migration forces. The observed period ratio distribution is inconsistent with either of those two scenarios on its own.  However, applying both stochastic and smooth migration forces to the planets simultaneously results in a period ratio distribution that is similar to the observed one. This is a natural \\rev{scenario} if planets form in a turbulent proto-planetary disk where these forces are always present. We show how the observed period ratio and eccentricity distribution can constrain the relative strength of these forces, a parameter which has been notoriously hard to predict for decades.  We make the source code \\rev{of our simulations} and the initial conditions freely available to enable the community to expand this study and include effect other than planetary migration. ",
        "introduction": "The number of discovered extrasolar planets is increasing rapidly. At the time when this letter was submitted, the number of planets has reached 784\\footnote{\\url{http://exoplanet.hanno-rein.de}}. Of those, 275 \\rev{(35\\%)} are in multi-planetary systems with two or more planets orbiting the same star. These systems are of particular interest to theorists as they can provide valuable information about their formation history.  The existence of mean motion resonances (MMRs) has been confirmed in multiple systems. The most studied planetary system in a MMR is Gliese~876. It consists of three gas giants which are locked in a tight 1:2:4 Laplace resonance \\rev{\\citep{Rivera2010,MarcyButler01}}. A large number of studies \\citep[e.g.][]{LeePeale01,LeePeale2002,SnellgrovePapaloizouNelson01,NelsonPapaloizou2002,beamic2003,veras2007} suggest that migration driven by a variety of mechanisms has played an important role in shaping this system.  This is not surprising as planet migration is a natural outcome of the interaction of a planet with the proto-planetary disk that it forms in \\citep{GoldreichTremaine1980}. Both $N$-body and hydrodynamical models can easily reproduce the observed period ratio, the eccentricities and the libration pattern even though the precise speed at which migration occurs is still up for debate. Furthermore, it is possible to place limits on the strength of additional stochastic forces which might be present in a turbulent proto-planetary disk \\rev{\\citep[e.g.][]{ReinPapaloizou2009}}.  Most of these planets have been discovered with the radial velocity method. This method is biased towards finding heavy planets on close-in orbits. Since 2009, the Kepler spacecraft is monitoring over one hundred thousand stars. Kepler has discovered thousands of planet candidates which now await confirmation \\citep{Batalha2012}. These planets have on average a much smaller mass than those discovered by radial velocity. Kepler therefore opens a new window to test planet formation scenarios.  In this letter, we apply a model of smooth migration (i.e. Type~I or Type~II depending on disk and planet properties) as well as stochastic migration forces to each candidate system with multiple planets. We show that even a small amount of \\rev{smooth} migration produces a distribution of period ratio that is inconsistent with observations. At each resonance, planets pile up, leading to distinct features in the cumulative distribution function of period ratios.  When stochastic migration forces are added, these features are smeared out. By adding just the right amount, one can retain some of the features, leading to a period ratio distribution similar to the observed one. Our simulations show a pile-up just outside of the exact commensurability. This is also seen in the Kepler data. However, it is less apparent in our simulations. This could either be resolved by fine-tuning several parameters or by adding additional physics such as tidal damping or gap edges. ",
        "conclusions": "\\label{sec:discussion} \\begin{figure} \\centering \\resizebox{0.99\\columnwidth}{!}{\\includegraphics{cdf_innerperiod}} \\caption{\\rev{Cumulative distribution of the period ratios in KOI systems divided into two bins: planetary systems where the period of the innermost planet is shorter/longer than 5 days. \\label{fig:cdf_innerperiod}}} \\end{figure}  Our results show that neither smooth nor stochastic planetary migration alone can reproduce the observed period ratio distribution of multi-planetary systems in the Kepler sample.  However, a combination of those two effects can create a period ratio distribution which is similar to the observed one. If this scenario is true, then we can use the eccentricity distribution of Kepler planets to constrain the relative strength of stochastic and smooth migration forces.  Other ideas have been brought up to explain the observed period ratios. \\cite{TerquemPapaloizou2007} show that migration of multiple planets in a disk with an inner edge, together with orbital circularization causes strict commensurability to be lost. \\rev{Similar scenarios involving tidal interactions have been studied by \\cite{Delisle2012,BatyginMorbidelli2012,LithwickWu2012}.} There is one important difference to the migration scenario presented here. The inclusion of tides leads a strong radial dependence of the effect. Such a dependency is currently not observed in the Kepler data set. \\rev{This can be verified in Figure~\\ref{fig:cdf_innerperiod} where the cumulative distribution is divided into two bins with almost equal number of planets. The two bins contain systems where the innermost planet's period is shorter/longer than 5 days. If tides were an important factor in shaping this distribution, one would expect a bigger effect for the bin that contains close-in planets. However, both distributions are identical. }  Although this letter focused on the migration scenario, we are interested in seeing other ideas being tested in a framework similar to the one presented in this letter. We therefore decided to make the initial conditions, the integrator and all plotting routines freely available. They can be downloaded as a tar-file from \\url{https://github.com/hannorein/rebound/tree/resonancelocation}. The relevant files for this project are located in the directory \\texttt{problems/resonancelocation}. We hope that this enables the community to further investigate these extremely interesting planetary systems and come up with new and maybe even better ideas."
    },
    "1208/1208.0828_arXiv.txt": {
        "abstract": "Very recently, it was pointed out that there exists a population of gamma-ray sources without associations at other wavelengths which exhibit spectral features consistent with mono-energetic lines at energies of approximately 111 and 129 GeV. Given recent evidence of similar gamma-ray lines from the Inner Galaxy, it is tempting to interpret these unassociated sources as nearby dark matter subhalos, powered by ongoing annihilations. In this paper, we study the spectrum, luminosity, and angular distribution of these sources, with the intention of testing the hypothesis that they are, in fact, dark matter subhalos. We find that of the 12 sources containing at least one prospective line photon, only 2 exhibit an overall gamma-ray spectrum which is consistent with that predicted from dark matter annihilations (2FGL J2351.6-7558 and 2FGL J0555.9-4348). After discounting the 10 clearly non-dark matter sources, the statistical significance of the remaining two prospective line photons is negligible. That being said, we cannot rule out the possibility that either or both of these sources are dark matter subhalos; their overall luminosity and galactic latitude distribution are not inconsistent with a dark matter origin. ",
        "introduction": "In many models, dark matter particles can annihilate through loop-level diagrams to final states which include a photon, such as $\\gamma \\gamma$, $\\gamma Z$ or $\\gamma h$. As dark matter annihilations in the universe today occur at non-relativistic velocities, the photons produced through such processes are predicted to take the form of mono-energetic gamma-ray lines. The prospect of observing such a gamma-ray line has long been considered a ``smoking gun'' for dark matter's indirect detection, providing a distinctive signal which is unlikely to be mimicked by astrophysical backgrounds.  The Fermi Gamma-Ray Space Telescope (Fermi-LAT)~\\citep{atwood_fermi} is one of the most promising instruments with which to detect such a gamma-ray line. Over the past few months, a great deal of attention has been given to the possibility that evidence for such gamma-ray lines is, in fact, present within the publicly available data from the Fermi-LAT~\\cite{Bringmann:2012vr,Weniger:2012tx}. The primary line in question appears at an energy of approximately 129 GeV, along with less statistically significant hints of a second line at around 111 GeV~\\cite{Su:2012ft}. These photons have been detected primarily within regions of the sky which reside within $\\sim$10$^{\\circ}$ of the Galactic Center, resembling the distribution predicted from a cusped halo profile. For recent discussions of these possible signals and their implications for particle physics, see Refs.~\\cite{Su:2012ft,data,dougidm,Boyarsky:2012ca} and~\\cite{pp,Buckley:2012ws}, respectively.  The task now at hand for the particle-astrophysics community is to determine whether the gamma-ray line or lines observed from the Inner Galaxy are 1) actually the products of dark matter annihilations, 2) somehow the result of a combination of astrophysical sources or mechanisms~\\cite{bubble,Aharonian:2012cs}, or 3) the result of systematic or instrumental issues~\\cite{Su:2012ft,dougidm,Boyarsky:2012ca}, perhaps associated with the inner workings of the Fermi-LAT itself. At this time, it is difficult to take very seriously the possibility that the appearance of the 129 GeV line is a statistical fluctuation, in particular in light of the very high significance of the line as derived from the template analysis of Ref.~\\cite{Su:2012ft} (greater than 5$\\sigma$, after accounting for an appropriate trials factor).  One way to possibly strengthen the case for a dark matter origin of the observed line would be to observe a line at the same energy from other directions of the sky thought to contain significant densities of dark matter. Such a signal has been recently reported, for example, from a number of Galaxy Clusters, although with only modest statistical significance~\\cite{Hektor:2012kc}. Even more recently, evidence of gamma-rays lines at both 129 and 111 GeV from a collection of unassociated gamma-ray sources has been presented~\\cite{new}. If confirmed, this result could have great bearing on the question of the line's origin. The Milky Way is predicted to contain large numbers of smaller subhalos and nearby subhalos could plausibly appear as a population of gamma-ray sources without counterparts at other wavelengths~\\cite{Belikov:2011pu,Mirabal:2012em,Buckley:2010vg,Zechlin:2011kk,acts,Kuhlen:2008aw,Zechlin:2011wa,Mirabal:2010ny}. If the line emission from such sources is confirmed, this would be most easily interpreted as evidence of dark matter annihilations taking place within such objects.  In this paper, we study the gamma-ray spectrum, luminosity, and sky distribution of these prospective line-emitting sources with the goal of assessing the likelihood that they are in fact dark matter subhalos. If a sizable fraction of these sources were shown to be subhalo-like in these regards, it would help to strengthen the case that the gamma-ray lines reported over the past several months do, in fact, originate from dark matter annihilations. What we actually find, however, is that the majority of these objects do not appear to be dark matter subhalos. That fact that line emission is observed from these objects, which do not appear to be powered by dark matter annihilations, adds credibility to the hypothesis that the lines reported from the Inner Galaxy are the product of some yet unknown systematic or instrumental issue, possibly associated with the Fermi-LAT (see also Ref.~\\cite{Su:2012ft,dougidm}). ",
        "conclusions": "In this paper, we have studied the spectra, luminosities, and sky distribution of the 12 unassociated 2FGL gamma-ray sources which exhibit possible line-emission at 129 or 111 GeV, as discussed in the recent Ref.~\\cite{new}. The conclusions of our study can be summarized as follows: \\begin{itemize} \\item{The majority of these 12 sources are very unlikely to be dark matter subhalos. In particular, 10 out of the 12 sources have measured spectra which are inconsistent with that predicted from dark matter annihilations. Furthermore, these sources are overwhelmingly concentrated at low galactic latitudes, in sharp contrast to that expected from a population of dark matter subhalos.} \\item{If we consider only the two 2FGL sources which exhibit spectra that are consistent with dark matter annihilation continuum emission (J0555.9-4348 and J2351.6-7558), the possibility that either or both of these sources may be subhalos cannot be ruled out. These two sources reside at moderate galactic latitudes (-40.6$^\\circ$ and -27.8$^\\circ$) and have overall fluxes which are consistent with the constraints on gamma-ray continuum emission.} \\end{itemize}  Looking at this problem from another perspective, the first of these two points implies that most (at least 11 out of the total 13) of the prospective line photons from unassociated 2FGL sources, as identified in Ref.~\\cite{new}, do not likely originate from dark matter annihilations, and instead require some other explanation. As gamma-ray lines of uniform energy are not observed or predicted from any known class of astrophysical objects, the possibility of an astrophysical solution seems unlikely. Instead, the observation of a line-like-signal from these unassociated sources strengthens the case that the reports of gamma-rays lines over the past several months are likely to be the result of some as of yet not understood systematic effect or instrumental issue.    This hypothesis will be put to the test by future observations by other gamma-ray telescopes. Specifically, pointed observations from HESS-II~\\cite{Bergstrom:2012vd} will be instrumental in determining whether the observed line signals are real or not. Additionally, the Gamma-400 instrument promises a significant improvement in both the angular and energy resolution, which will enable a detailed comparison between the tentative line feature and the continuum emission predicted in dark matter scenarios~\\citep{gamma400}.     \\bigskip   {\\it Acknowledgements:} We would like to thank Doug Finkbeiner for valuable discussions.  DH is supported by the US Department of Energy, including grant DE-FG02-95ER40896."
    },
    "1208/1208.0533_arXiv.txt": {
        "abstract": "We calculate merger rates of dark matter haloes using the Extended Press-Schechter approximation (EPS) for the Spherical Collapse (SC) and the Ellipsoidal Collapse (EC) models.\\\\ Merger rates have been calculated for masses in the range $10^{10}M_{\\odot}\\mathrm{h}^{-1}$ to $10^{14}M_{\\odot}\\mathrm{h}^{-1}$ and for redshifts  $z$ in the range $0$ to $3$ and they have been compared with merger rates that have been proposed by other authors as fits to the results of N-body simulations. The detailed comparison presented here shows that the agreement between the analytical models and N-body simulations depends crucially on the mass of the descendant halo. For some range of masses and redshifts either SC or EC models approximate satisfactory the results of N-body simulations but for other cases both models are less satisfactory or even bad approximations. We showed, by studying the parameters of the problem that a disagreement --if it appears--  does not depend on the values of the parameters but on the kind of the particular solution used for the distribution of progenitors or on the nature of EPS methods.\\\\ Further studies could  help to improve our understanding about the physical processes during the formation of dark matter haloes. ",
        "introduction": "The development of analytical or semi-numerical methods for the problem of structure formation in the universe helps to improve our understanding of important physical processes. A class of such methods is based on the ideas of \\citet{prsc74} and on their extensions (Extended Press-Schechter Methods EPS,  \\citet{boet91}, \\citet{laco93}): The linear overdensity $\\delta(\\textbf{x};R)\\equiv [\\rho(\\textbf{x};R)-\\rho_b(\\textbf{x};R)]/\\rho_b(\\textbf{x};R)$ at a given point $\\textbf{x}$ of an initial snapshot of the Universe fluctuates when the smoothing scale $R$ decreases. In the above relation, $\\rho(\\textbf{x};R)$ is the density at point $\\textbf{x}$ of the initial Universe smoothed by a window function with smoothing scale $R$. The index $b$ denotes the density of the background model of the Universe. This fluctuation is a Markovian process when the smoothing is performed using a top-hat window in Fourier space. For any value of the smoothing scale $R$, the overdensity field is assumed to be Gaussian with zero mean value. The dispersion of these Gaussians is a decreasing function of the smoothing scale $R$ reflecting the large scale homogeneity of the Universe. The mass $M$ contained in a given scale $R$ depends on the window function used. For the top-hat window this relation is: $M=\\frac{4}{3}\\pi\\rho_{b,i}R^3=\\frac{\\Omega_{m,i}H^2_i}{2G}R^3$, where $\\rho_{b,i}$ and $\\Omega_{m,i}$ are the values of the mean density and the density parameter of the Universe, $G$ is the gravitational constant and $H_i$ is the Hubble's constant. The index $i$ indicates that all the above values are calculated at the initial snapshot. The dispersion in mass,$\\sigma^2$, at scale $R$ is a function of mass $M$ and it is usually denoted by $S$, that is $S(M)\\equiv\\sigma^2[R(M)]$.\\\\ In the plane $(S,\\delta)$ random walks start from the point $(S=0,\\delta=0)$ and diffuse as $S$ increases. Let the line $B=B_{SC}(z)$ that is a function of redshift $z$. In the case this line is parallel to $S$-axis in the $(S,\\delta)$ plane, then  it has a physical meaning as it can be connected to the spherical collapse model (SC): It is well known that in an Einstein-de Sitter Universe, a spherical overdensity collapses at $z$ if the linear extrapolation of its value up to the present exceeds $\\delta_{sc}\\approx 1.686$ (see for example \\citet{peeb80}). All involved quantities (density, overdensities, dispersion) are linearly extrapolated to the present and thus the barrier in the spherical collapse model is written in the from $B(z)=1.686/D(z)$, where $D(z)$ is the growth factor derived by the linear theory, normalized to unity at the present epoch. It is clear that the line $B(z)$ is an increasing function of $z$. If a random walk crosses this barrier for first time at some value $S_0$ of $S$, then the mass element associated with the random walk is considered to belong to a halo of mass $M_0=S^{-1}(S_0)$ at the epoch with redshift $z$. However, the distribution of haloes mass $f_M$, at some epoch $z$, is connected  to the first crossing distribution $f_S$, by the random walks, of the barrier that corresponds to epoch $z$ with the relation: \\begin{equation} f_M(M)\\mathrm{d}M=f_S(S)|\\frac{\\mathrm{d}S(M)}{\\mathrm{d}M}|\\mathrm{d}M \\end{equation} A form of the barrier that results to a mass function that is in better agreement with the results of N-body simulations than the spherical model is the one given by the Eq. \\begin{equation} B_{EC}(S,z)=\\sqrt{a}B_{SC}(z)[1+\\beta[S/aB_{SC}^2(z)]^{\\gamma}]. \\end{equation} In the above Eq. $\\alpha$, $\\beta$ and $\\gamma$ are constants. The above barrier represents an ellipsoidal  collapse model (EC) \\citep{shto99}. The barrier  depends on the mass ($S=S(M)$) and it is called a moving barrier. The values of the parameters are $a=0.707$,~$\\beta=0.485$,~ $\\gamma=0.615$ and are adopted either from the dynamics of ellipsoidal collapse or from fits to the results of N-body simulations The spherical collapse model results for $a=1$ and $\\beta=0.$\\\\ In a hierarchical scenario of the formation of haloes, the following question is fundamental: Given that at some redshift $z_0$ a mass element belongs to a halo of mass $M_0$, what is the probability the same mass element at some larger redshift $z$ $(z >z_0)$ -that corresponds to an earlier time- was part of a halo with mass $M$ with $ M\\leq M_0$? This question in terms of first crossing distributions and barriers can be written in the following equivalent form:  Given  a random walk passes for the first time from the point $(\\delta_0,S_0)$ what is the probability this random walk crosses a barrier $B$ with $B > \\delta_0$, for the first time between $S,~S+\\mathrm{d}S$ with $S > S_0$?\\\\ If we denote the above probability by $f(S/\\delta_0,S_0)\\mathrm{d}S$ it can be proved, \\citep{zhhu06}, that for an arbitrary barrier, $f$ satisfies the following integral equation: \\begin{equation} f(S/\\delta_0,S_0)=g_1(S,\\delta_0,S_0)+\\int_0^Sg_2(S,S')f(S'/\\delta_0,S_0)\\mathrm{d}S' \\end{equation} where: \\begin{equation} g_1(S,\\delta_0,S_0)=\\left[\\frac{B(S)-\\delta_0}{S-S_0}-2\\frac{\\mathrm{d}B(S)}{\\mathrm{d}S}\\right]P_0[B(S),S/\\delta_0,S_0] \\end{equation} \\begin{equation} g_2(S,S')=\\left[2\\frac{\\mathrm{d}B(S)}{\\mathrm{d}S}-\\frac{B(S)-B(S')}{S-S'}\\right]P_0[B(S),S/B(S'),S'] \\end{equation} and \\begin{equation} P_0(x,y/x_0,y_0)=\\frac{1}{\\sqrt{2\\pi(\\Delta y)}}e^{- \\frac{{\\Delta x}^2}{2\\Delta y}} \\end{equation} with $\\Delta y=x-x_0$ and $\\Delta y=y-y_0$.\\\\ In the case of a linear barrier Eq.(3) admits an analytic solution. If $B(S)=\\omega+qS$, where the coefficients $\\omega$ and $q$ could be functions of the redshift $z$ in order to describe the dependence on the time, the solution is written: \\begin{equation} f(S /\\delta_0,S_0)=\\frac{B(S_0)-\\delta_0}{\\sqrt{2\\pi(S-S_0)^3} }\\exp\\left[-\\frac{[B(S)- \\delta_0]^2}{2(S-S_0)}\\right] \\end{equation} Thus, the spherical model which is of the form $B=B(z)=\\omega(z)=1.686/D(z)$ leads to the solution: \\begin{equation} f_{SC}(S,z/S_0,z_0)=\\frac{\\Delta \\omega}{\\sqrt{2\\pi(\\Delta S)^3}}\\exp\\left[-\\frac{(\\Delta \\omega)^2}{2\\Delta S}\\right] \\end{equation} where, $\\Delta S\\equiv S-S_0$, and $\\Delta \\omega=\\omega(z)-\\omega(z_0)$ .\\\\ Unfortunately, no analytical solution exists for the ellipsoidal model. The exact numerical solution of Eq.(3) is well approximated by the expression proposed by \\citet{shto02} that is: \\begin{equation} f_{EC}(S,z/S_0,z_0)=\\frac{1}{\\sqrt{2\\pi}}\\frac{|T(S,z/S_0,z_0)|}{(\\Delta S)^{3/2}} \\exp\\left[-\\frac{(\\Delta B)^2}{2\\Delta S}\\right]\\mathrm{d}S \\end{equation} where, $\\Delta B=B(S,z)-B(S_0,z_0)$, and the function $T$ is given by: \\begin{equation} T(S,z/S_0,z_0)=B(S,z)-B(S_0,z_0)+\\sum_{n=1}^{5}\\frac{[S_0-S]^n}{n!} \\frac{\\mathrm{\\partial ^n}}{\\partial S^n}B(S,z). \\end{equation} According to the hierarchical clustering any halo is formed by smaller haloes (progenitors). A number of progenitors merge at $z$ and form a larger halo of mass $M_0$ at $z_0$ ($z_0 <z$). Obviously, the sum of the masses of the progenitors equals to $M_0$. Given a halo of mass $M_0$ at $z_0$  the average number of its progenitors in the mass interval $[M,M+\\mathrm{d}M]$ present at $z$ with $z > z_0$ is : \\begin{equation} \\frac{\\mathrm{d}N}{\\mathrm{d}M}(M/M_0,\\Delta \\omega)\\mathrm{d}M=\\frac{M_0}{M}f(S,z/S_0,z_0)\\mathrm{d}M \\end{equation} Recent comparisons  show that the use of EC model improves the agreement between the results of EPS methods and those of N-body simulations. For example, \\citet{yaet04} showed that the multiplicity function resulting from  N-body simulations is far from the predictions of spherical model while it shows an excellent agreement with the results of the EC model. On the other hand, \\citet{liet03} compared the distribution of formation times of haloes formed in N-body simulations with the formation times of haloes formed in terms of the spherical collapse model of the EPS theory. They found that N-body simulations give smaller formation times. \\citet{hipo06} showed that using the EC model, formation times are shifted to smaller values than those predicted by a spherical collapse model. Additionally, the EC model combined with the stable ``clustering hypothesis\" has been  used by \\citep{hi06} in order to study density profiles of dark matter haloes. Interesting enough,. the resulting density profiles at the central regions are closer to the results of observations than are the results of N-body simulations. Consequently, the EC model is a significant improvement of the spherical model and therefore we are  well motivated to study  merger-rates of dark matter haloes for both the SC and the EC model. This study depends upon  the accurate  construction of a set of progenitors for any halo for a very small ``time step\"  $\\Delta \\omega$. The set of progenitors are created using the method proposed by \\citet{nede08} that we describe in Sect.3. In Sect. 2 we define merger rates and we recall fitting formulae resulting from N-body simulations. In Sect.4 our results are presented and discussed. ",
        "conclusions": ""
    },
    "1208/1208.5911_arXiv.txt": {
        "abstract": "{About half  of the baryons of the Universe are  expected to be in the form of filaments of hot and low-density intergalactic medium. Most of these baryons remain undetected even by the most advanced X-ray observatories, which are limited in sensitivity to the diffuse low-density medium. } {The \\Planck\\ satellite has provided hundreds of detections of the hot gas in clusters of galaxies via the thermal Sunyaev-Zel'dovich (tSZ) effect and is an ideal instrument for studying extended low-density media through the tSZ effect. In this paper we use the \\Planck\\ data to search for signatures of a fraction of these missing baryons between pairs of galaxy clusters.} {Cluster pairs are good candidates for searching for the hotter and denser phase of the intergalactic medium (which is more easily observed through the SZ effect). Using an X-ray catalogue of clusters and the Planck data, we selected  physical pairs of clusters as candidates. Using the Planck data, we constructed a local map of the tSZ effect centred on each pair of galaxy clusters. \\ROSAT\\ data were used to construct X-ray maps of these pairs. After hmodelling and subtracting the tSZ effect and X-ray emission for each cluster in the pair, we studied the residuals on both the SZ and X-ray maps.} {For the merging cluster pair A399-A401 we observe a significant tSZ effect signal in the intercluster region beyond the virial radii of the clusters. A joint X-ray  SZ analysis allows us to constrain the temperature and density of this intercluster medium. We obtain a temperature of  $kT = 7.1 \\pm 0.9$ \\, keV (consistent with previous estimates) and a baryon density of $(3.7 \\pm 0.2) \\times 10^{-4}$ \\, cm$^{-3}$.} {The \\Planck\\  satellite mission has provided the first SZ detection of the hot and diffuse intercluster gas.} ",
        "introduction": "\\label{introduction}  A sizeable fraction of the baryons of the Universe are expected to be in the form of the WHIM (warm-hot intergalactic medium) and remain undetected at low redshifts. The WHIM is expected to exist mostly in filaments but also around and between massive clusters. These missing baryons are supposed to be in a low-density, low-temperature phase (overdensities between 5 to 200 times the critical density and $T = 10^5-10^7$ \\, K, \\cite{Cen99}), making the amount of X-rays produced by the WHIM too small to be detected with current X-ray facilities. By contrast, their detection could be possible via the Sunyaev-Zel'dovich (SZ hereafter) effect \\citep{SZ} produced by the inverse Compton scattering between the cosmic microwave background (CMB) photons  and the electrons of the WHIM. As the SZ effect is proportional to the electron pressure in the medium, low-density and low-temperature regions can be detected provided their integrated signal is strong enough. \\Planck's relatively poor resolution becomes an advantage in this situation since it permits scanning of wide regions of the sky that can later be integrated to increase the signal-to-noise ratio of the diffuse (but intrinsically large-scale) SZ signal.  The full-sky coverage and wide frequency range of the \\Planck\\ satellite mission makes it possible to produce reliable maps of the tSZ emission \\citep{2011A&A...536A...8P,2011A&A...536A...9P,2011A&A...536A..10P}. In particular, Planck is better suited than ground experiments to detecting diffuse SZ signals, such as the WHIM, which can extend over relatively large angular scales. Ground experiments can be affected at large angular scales by atmospheric fluctuations that need to be removed. This removal process can distort the modes that include the large angular scale signals. \\Planck\\ data do not suffer from these limitations and can use their relatively poor angular resolution (when compared to some ground experiments) to its advantage. Indeed, diffuse low surface brightness objects can be resolved and detected by \\Planck. Finally, the wide frequency coverage and extremely high sensitivity of \\Planck\\ allows for detailed foreground (and CMB) removal that otherwise would overwhelm the weak signal of the WHIM.  The gas around clusters is expected to be hotter and denser than the WHIM in filaments, making direct detection of the cluster gas more likely. In addition, the increase of pressure caused by the merging process enhances the SZ signal, making it easier to detect the gas between pairs of interacting clusters. In the process of hierarchical formation clusters assemble via continuous accretion and merger events. Therefore, the bridge of intercluster matter between them is expected to be of higher density, temperature, and thus thermal pressure than the average WHIM matter found in cosmic filaments \\citep{dolag2006}. \\\\  The \\Planck\\  satellite (\\cite{planckmission}) has the potential to detect these filamentary structures directly via the SZ effect. Suitable targets for Planck are close objects that subtend large solid angles and therefore have high integrated SZ fluxes. Alternatively, regions between mergers (filaments between pairs of clusters) or extremely deep gravitational wells (superclusters such as the Shapley or Corona Borealis, \\cite{cacho2009}) will contain diffuse gas with increased pressure that could be detected by \\Planck.  For this work, we concentrate on searching for diffuse filamentary-like structure between pairs of merging clusters. We used the MCXC (Meta-Catalog of X-Ray Detected Clusters) catalogue of clusters of galaxies \\citep{Piffaretti2011} and the \\Planck\\ data to select a sample of pairs of merging clusters to study the properties of the gas in the intercluster region. \\\\  Indirect WHIM detections have been claimed through absorption lines in the X-ray (and UV) band  \\citep{Richter08}. There is also evidence of filamentary structure in the intercluster region from X-ray observations of several well-known merging cluster pairs such as A222-A223 \\citep{Werner2008}, A399-A401 \\citep{Sakelliou2004}, A3391-A3395 \\citep{Tittley2001},  and from the double cluster A1758 \\citep{Durret}. The pairs of clusters A3391-A3395 (separated by about 50$\\arcm$ on the sky and at redshifts z=0.051 and z=0.057, respectively \\citep{Tittley2001}) and more specially, A399-A401 (separated by about 40$\\arcm$ on the sky and at redshifts  z=0.0724  and z=0.0737, respectively) are of particular interest for the purpose of this paper, given their geometry and angular separation. This is sufficient to allow \\Planck\\ to resolve the individual cluster components.  For A399-A401, earlier observations show an excess of X-ray emission above the background level in the intercluster region. Using XMM data, \\cite{Sakelliou2004} obtained  best-fitting models in the intercluster region that indicated such an excess. Both clusters are classified as non-cool-core clusters and show weak radio halos \\citep{Murgia2010}. These two facts could be an indication of a past interaction between the two clusters. \\cite{Fujita1996}  analysed  \\ASCA\\ data of the intercluster region and found a relatively high temperature in this region. They suggested a pre-merger scenario but did not rule out a past interaction. \\cite{Fabian1997} used HRI \\ROSAT\\ data and found a prominent linear feature in A399 pointing towards A401. They suggested that this could be evidence of a past interaction. Using \\Suzaku\\  observations, \\cite{Fujita2008} found that the intercluster region has a relatively high metallicity of 0.2 solar. These works estimated that the filamentary bridge has an electron density of $n_e \\sim 10^{-4}$ \\, cm$^{-3}$ \\citep{Fujita1996,Sakelliou2004,Fujita2008}.  \\\\   In this paper we concentrate on pairs of merging clusters including A399-A401 and we study the physical properties of the gas in the intercluster region via a combined analysis of the  tSZ effect and the X-ray emission. The paper is organized as follows. Section \\ref{P_Data} gives a brief description of the \\Planck\\ data used for this study. In Section~\\ref{selection} we describe the selection procedure used to identify the most suitable pairs of clusters for the analysis. We search for pairs of clusters for which the contribution of the SZ effect to the signal is significant in the intercluster medium. Section~\\ref{xrays} describes the X-ray \\ROSAT\\ observations for the selected pairs of clusters. In section~\\ref{Models} we model the SZ and X-ray emission from the clusters assuming spherical symmetry and subtract them from the data. In section~\\ref{analysis} the SZ and X-ray residuals are fitted to a simplified filament model to characterize the physical properties of the intercluster region. Section~\\ref{discussions} discusses our main results focusing on the limitations imposed by the cluster spherical symmetry assumption and alternative non-symmetric scenarios. Finally, we conclude in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  Using \\Planck\\ and \\ROSAT\\ data, we have studied the tSZ and X-ray maps of 25 pairs of clusters of galaxies. After modelling (assuming a spherical symmetric model) and subtracting the contribution of each individual cluster, we detected significant tSZ residuals in at least two of these pairs: A399-A401 and A3391-A3395. In the case of the A399-A401 pair, these residuals are compatible with an intercluster filament of hot, 7 keV (in agreement with \\cite{Sakelliou2004}), and diffuse, $3.7 \\times 10^{-4}$ \\,cm$^{-3}$, gas connecting the two clusters. A chance coincidence of a background cluster is ruled out for canonical scaling relations because it would have to be a very massive cluster ($M_{500} = 2.4 \\times 10^{15} \\, M_{\\odot}$) at a very high redshift ($z=1.9$).  The signal detected by \\Planck\\ is significant independently of the cluster model. Hydrodynamical simulations show that the intercluster signal in A399-A401 is  compatible with a scenario where the intercluster region is populated with a mixture of material from the clusters and the intergalactic medium, indicating that there might be a bridge of matter connecting the two clusters. If the measured signal of the merging cluster pair can be interpreted in terms of spherically symmetric individual clusters, evidence remains for an intercluster SZ signal detected by \\Planck. It is consistent with simulated data and may constitute the first detection of the tSZ effect between clusters. Under this interpretation, the signal is unambiguous in the sense that it is detected with high significance (as shown by Figure 2), is not caused by any known artefact, is clearly resolved by Planck, and is located in what one would consider the external region of a standard cluster. The exact interpretation of the origin of the signal is more open to speculation. The analysis presented in this work shows the potential of \\Planck\\ data for studying these yet unexplored regions. Better angular resolution observations of the tSZ would improve the modelling of the  clusters and reduce the uncertainties in the estimation of the signal excess in the intercluster region."
    },
    "1208/1208.3706_arXiv.txt": {
        "abstract": "We present an analysis of the classic Alcubierre metric based on conformal gravity, rather than standard general relativity. The main characteristics of the resulting warp drive remain the same as in the original study by Alcubierre, namely that effective super-luminal motion is a viable outcome of the metric.  We show that for particular choices of the shaping function, the Alcubierre metric in the context of conformal gravity does not violate the weak energy condition, as was the case of the original solution. In particular, the resulting warp drive does not require the use of exotic matter.  Therefore, if conformal gravity is a correct extension of general relativity, super-luminal motion via an Alcubierre metric might be a realistic solution, thus allowing faster-than-light interstellar travel. ",
        "introduction": "Introduction}  In 1994 M. Alcubierre introduced the so-called \\textit{Warp Drive Metric }(WDM), within the framework of General Relativity (GR), which allows in principle for super-luminal motion, i.e., faster-than-light travel \\cite{Alcubierre:1994tu}. This super-luminal propulsion is achieved by respectively expanding and contracting the space-time behind and in front of a spaceship, while the spacecraft is left inside a locally flat region of space-time, within the so-called \\textit{warp bubble}.  In this way the spaceship can travel at arbitrarily high speeds, without violating the laws of special and general relativity, or other known physical laws. Furthermore, the spacecraft and its occupants would also be at rest in flat space-time, thus immune from high accelerations and unaffected by special relativistic effects, such as time dilation. Enormous tidal forces would only be present near the edge of the warp bubble, which can be made large enough to accommodate the volume occupied by the ship.  However, Alcubierre \\cite{Alcubierre:1994tu}\\ was also the first to point out that this hypothetical solution of Einstein's equations of GR would violate all three standard energy conditions (weak, dominant and strong; see \\cite{Hawking1}, \\cite{Carroll:2004st}, and \\cite{Wald1}\\ for definitions). In particular, the violation of the weak energy condition (WEC) implies that negative energy density is required to establish the Alcubierre WDM, thus making it practically impossible to achieve this type of super-luminal motion, unless large quantities of exotic matter (i.e., with negative energy density) can be created. Since our current knowledge of this type of exotic matter is limited to some special effects in quantum field theory (such as the Casimir effect), it is unlikely that the Alcubierre WDM can be practically established within the framework of General Relativity.  Following Alcubierre's seminal paper, many other studies appeared in the literature, either proposing alternatives to the original warp drive mechanism (\\cite{Krasnikov:1995ad}, \\cite{Natario:2001tk}) or refining and analyzing in more detail the original idea (\\cite{Olum:1998mu}, \\cite{VanDenBroeck:1999sn}, \\cite{Clark:1999zn}, \\cite{GonzalezDiaz:1999db}, \\cite{Pfenning:1997wh}, \\cite{2003GReGr..35.2025W}, \\cite{Lobo:2004wq}, \\cite{GonzalezDiaz:2007zza}, \\cite{GonzalezDiaz:2007zzb}, \\cite{Finazzi:2009jb}, \\cite{Barcelo:2010pu}, \\cite{Muller:2011fa}, \\cite{McMonigal:2012ey}). However, all these studies were conducted using standard GR and could not avoid the violation of the WEC, meaning that some exotic matter would always be required for faster-than-light travel. Similar issues also exist in other well-known GR solutions for super-luminal motion, such as space-time wormholes \\cite{Hartle1}.  Einstein's General Relativity and the related \\textquotedblleft Standard Model\\textquotedblright\\ of Cosmology have been highly successful in describing our Universe, from the Solar System up to the largest cosmological scales, but recently these theories have also led to a profound crisis in our understanding of its ultimate composition. From the original discovery of the expansion of the Universe, which resulted in standard Big Bang Cosmology, scientists have progressed a long way towards our current picture, in which the contents of the Universe are today described in terms of two main components, dark matter (DM) and dark energy (DE), accounting for most of the observed Universe, with ordinary matter just playing a minor role.  Since there is no evidence available yet as to the real nature of dark matter and dark energy, alternative gravitational and cosmological theories are being developed, in addition to standard explanations of dark matter/dark energy invoking the existence of exotic new particles also yet to be discovered. In line with these possible new theoretical ideas, \\textit{Conformal Gravity} (CG) has emerged as a non-standard extension of Einstein's GR, based on a possible symmetry of the Universe: the conformal symmetry, i.e., the invariability of the space-time fabric under local \\textquotedblleft stretching\\textquotedblright\\ of the metric (for reviews see \\cite{Mannheim:2005bfa}, \\cite{Varieschi:2008fc}). This alternative theory has been re-introduced in recent years (following the original work by H. Weyl \\cite{Weyl:1918aa}, \\cite{Weyl:1918ib}, \\cite{Weyl:1919fi}), leading to cosmological models which do not require the existence of DM and DE (\\cite{Mannheim:2011ds}, \\cite{Mannheim:2010ti}, \\cite{Mannheim:2010xw}, \\cite{O'Brien:2011wg}, \\cite{Mannheim:2007ug}, \\cite{Varieschi:2008va}, \\cite{Diaferio:2011kc}).  In view of a possible extension of Einstein's General Relativity into Conformal Gravity, in this paper we have re-considered the Alcubierre WDM, basing it on CG rather than standard GR. In Sect. \\ref{sect:conformal_gravity} we review the fundamental principles of CG and the calculation of the stress-energy tensor in this gravitational theory. In Sect. \\ref{sect:alcubierre_metric} we consider the Alcubierre metric in CG and compute the energy density for different shaping functions of the metric.  In particular, we will show that for certain shaping functions, the Alcubierre metric in the context of Conformal Gravity does not violate the weak energy condition, as was the case of the original solution. This analysis continues in Sect. \\ref{sect:energy_conditions}, where we study other energy conditions and estimate the total energy required for this CG warp drive. Finally, in Sect. \\ref{sect:conclusions}, we conclude that if CG is a correct extension of GR, super-luminal motion via an Alcubierre metric might be a realistic possibility, thus enabling faster-than-light interstellar travel without requiring exotic matter. ",
        "conclusions": "Conclusions}  In this paper we have analyzed in detail the Alcubierre warp drive mechanism within the framework of Conformal Gravity. We have seen that a particular choice of the shaping function (Hartle shaping function, instead of the original Alcubierre one) can overcome the main limitation of the AWD in standard General Relativity, namely the violation of the weak energy condition.  In fact, we have shown that for a wide range of spaceship velocities, the CG solutions do not violate the WEC, and, therefore, the AWD mechanism might be viable, if CG is the correct extension of the current gravitational theories. All the components of the stress-energy tensor can be analytically calculated, using a Mathematica program based on Conformal Gravity. Thus, a warp drive can, at least in principle, be fully established following our computations.  We have also checked two other main energy conditions: the SEC is always verified, while the DEC is violated, at least in the case we considered. Finally, we estimated the energy needed to establish a reasonable warp drive at the speed of light. This energy depends critically on the value of $\\alpha_{g}$, the CG coupling constant, which is not well known. Therefore, this estimate will need to be refined in future studies."
    },
    "1208/1208.4256_arXiv.txt": {
        "abstract": "The conversion of gas into stars is a fundamental process in astrophysics and cosmology.  Stars are known to form from the gravitational collapse of dense clumps in interstellar molecular clouds, and it has been proposed that the resulting star formation rate is proportional to either the amount of mass above a threshold gas surface density, or the gas volume density.  These star-formation prescriptions appear to hold in nearby molecular clouds in our Milky Way Galaxy's disk as well as in distant galaxies where the star formation rates are often much larger.  The inner 500\\,pc of our Galaxy, the Central Molecular Zone (CMZ), contains the largest concentration of dense, high-surface density molecular gas in the Milky Way, providing an environment where the validity of star-formation prescriptions can be tested.  Here we show that by several measures, the current star formation rate in the CMZ is an order-of-magnitude lower than the rates predicted by the currently accepted prescriptions.  In particular, the region $1^\\circ < l < 3.5^\\circ$, $|b| < 0.5^\\circ$ contains $\\sim10^7$\\,M$_\\odot$ of dense molecular gas --- enough to form 1000 Orion-like clusters --- but the present-day star formation rate within this gas is only equivalent to that in Orion. In addition to density, another property of molecular clouds, such as the amplitude of turbulent motions, must be included in the star-formation prescription to predict the star formation rate in a given mass of molecular gas. ",
        "introduction": "Stars play a pivotal role in shaping the cosmos. High-mass stars contributed to the ionisation of the Universe after the Cosmic Dark Ages.  They drive energy cycles and chemical enrichment on galactic scales, hence sculpting galactic structure. Thus, the conversion of gas into stars is fundamental to astrophysics and cosmology.  The rate at which gas is converted into stars has been measured in the disks of nearby galaxies.  When averaged over hundreds of parsecs, the star-formation rate (SFR) was found to have a power-law dependence on the gas surface-density as described by the Schmidt-Kennicutt (SK) relations \\citep{schmidt1959,kennicutt1998,kennicutt_evans2012}.  A linear relationship is found between SFR and gas surface-density above a local extinction threshold, A$_{\\rm K} \\sim$0.8\\, magnitudes at a near-infrared wavelength of 2.2 $\\mu$m, corresponding to a gas column-density of $\\sim7.4\\times10^{21}$ hydrogen molecules per cm$^{-2}$ or a gas surface-density $\\Sigma_{\\rm gas}$, of $\\sim116$\\,M$_\\odot$pc$^{-2}$ \\citep{lada2012} .  For a typical 0.1 pc radius cloud core, this corresponds to a volume density $ 3 \\times 10^4$\\, hydrogen molecules cm$^{-3}$.  Alternatively, it has been proposed that the SFR density is linearly proportional to the mean gas density divided by the free-fall time multiplied by an efficiency factor estimated to be about 1\\% \\citep{kdm2012}.  These surface- and volume-density relations potentially unify our understanding of SFRs from the nearest star forming regions to ultra-luminous infrared galaxies (ULIRGs) and even star forming regions at intermediate to high redshifts \\citep{swinbank2010,danielson2011}.  With the aim of testing these star formation (SF) scaling relations across large volumes of the Milky Way (MW), in $\\S$~\\ref{sec:dg_sf_tracers} we start by looking at the choice of observational diagnostics of dense gas and star formation activity in the Galaxy. In $\\S$~\\ref{sec:obs_data}, we present the survey data used in this work. In $\\S$~\\ref{sec:discussion_sfr_per_unit_mass_in_gal}~ \\&~\\ref{sec:variation?} we discuss the striking differences between the dense gas and SF activity tracers across the Galaxy and investigate if any observational or systematic biases may be causing such differences. We find no evidence that the results are strongly affected by observational or systematic biases. In $\\S$~\\ref{sec:quant_analyis_cmz} we then perform a quantitative analysis of the SFR and gas mass in the CMZ and directly compare this to predictions of proposed column density threshold and volumetric SF relations. In $\\S$~\\ref{sec:summary} we summarise the results and discuss the implications for the universality of star formation relations. ",
        "conclusions": "As the Galaxy is optically-thin to the $\\nhone$, 500\\,$\\mu$m, water maser, methanol maser and RRL emission, the emission surface density ratios between these tracers in the plane of the sky is equivalent to that in a face-on view of the Galaxy. Comparing the emission ratios as a function of longitude and making the reasonable assumptions that i) the integrated intensity of the far-IR and $\\nhone$ emission is proportional to the mass of dense gas, and, ii) the number of water masers, methanol masers, \\hii\\ regions and RRL integrated intensity is proportional to the amount of SF, we now investigate the relationship between the dense gas and SFR in the Galaxy.  Figure~\\ref{fig:long_nh3_mas} shows the dense gas emission and current star formation activity tracers as a function of Galactic longitude ($l$) and latitude ($b$).  The dense gas distribution is dominated by the very bright and spatially-extended emission within a few degrees of the Galactic centre -- the CMZ \\citep{morris_serabyn1996, ferriere2007}. This is easily distinguished from the rest of the Galactic plane by the very high intensity of dense gas emission and large gas velocity dispersion as indicated by wide spectral lines. While the dense gas emission is highly concentrated in the CMZ, the distribution of star formation activity tracers is relatively uniform across the Galaxy. Quantitatively, the CMZ accounts for $\\sim$80\\% of the integrated $\\nhone$ intensity but only contains 4\\% of the survey area.  Yet, the CMZ does not stand out in either water or methanol maser emission, or in the number of \\hii~regions, which all trace recent high-mass star formation. A qualitatively similar trend is reported by \\citet{beuther2012} who compare the sub-mm dust continuum emission and GLIMPSE point sources as a function of Galactic longitude.  Figure~\\ref{fig:long_nh3_mas_hist} shows Galactic longitude distributions of dense gas and star-formation activity indicators summed over the observed latitude range ($-0.5^\\circ \\leq b \\leq 0.5^\\circ$) for each longitude pixel. To make a direct comparison of dense gas and SF tracers as a function of longitude, we first sampled in 2-degree longitude bins to ensure the volume of the Galaxy covered in each bin contains a large number of SF regions and so is appropriate for testing SF relations \\citep{onodera2010}. The emission or number count in each 2 degree longitude bin was then normalised by the total emission or number in the full survey area.  The top panel of Figure~\\ref{fig:dense_gas_vs_masers} shows that there is a strong correlation between independent dense gas tracers. To separate the CMZ emission from that in the rest of the Milky Way, from here on we refer to the region $|l|<5^\\circ$ with $\\nhone$ line widths $\\Delta$V$~>$~15\\,$\\kms$ as ``GC-only'', to distinguish it from the rest of the Galaxy, which we refer to as the ``non-GC'' region. No offset is seen in the correlation between the GC-only and non-GC regions. The middle panel of Figure~\\ref{fig:dense_gas_vs_masers} shows a correlation between the dense gas and SF tracers for the non-GC regions. However, the GC-only regions are clearly distinct, with at least an order of magnitude brighter dense gas emission for the number of SF activity tracers.  The bottom panel of Figure~\\ref{fig:dense_gas_vs_masers} shows the resulting dense gas vs SF tracer surface density ratio in Galacto-centric radius, R$_{\\rm GC}$, annuli of 0.5\\,kpc. R$_{\\rm GC}$ was calculated using the Galactic rotation curve of \\citet{brand_blitz1993} and assuming a distance to the Galactic centre of 8.5 kpc and a solar velocity of 220\\,kms$^{-1}$. The gas between the CMZ and R$_{\\rm GC} \\sim$3\\,kpc shows emission at anomalous velocities so the rotation curve does not place reliable constraints on R$_{\\rm GC}$ over this range. The surface density ratios over this region, highlighted by the hatched rectangle, should be ignored. The ratio is approximately constant at R$_{\\rm GC}$$>$3\\,kpc. This suggests the linear relationship observed between the quantity of gas above the proposed extinction threshold and the SFR \\citep{gao_solomon2004,wu2005,lada2012} extends to a larger number of more representative SF regions across the Galaxy. By comparison, the dense gas surface density towards the GC (R$_{\\rm GC}$$<$0.5\\,kpc) is orders of magnitude too large compared to the SF activity surface density.   \\label{sec:summary}  In summary, we find that the dense gas ($\\nhthree$ \\& 500\\,$\\mu$m) and SF activity tracers (masers \\& HII regions) used in this study are reliably tracing the present-day relative dense gas mass and SF activity distributions, respectively, across the Galaxy. We conclude that the striking difference between the dense gas and SF activity tracers between the GC-only and non-GC regions shows the current star formation rate per unit mass of dense gas is an order of magnitude smaller in the GC than in the rest of the Galaxy. We directly test the predictions of proposed column-density threshold and volumetric star formation relations and find, given the mass of dense gas in the GC, they over-predict the observed SFR in the GC by an order of magnitude. We conclude the current star formation relations are incomplete in some way. Any universal column/volume density relations must be a \\emph{necessary but not sufficient} condition for SF to occur.  Putting the CMZ in the context of the Galaxy as a whole, the Milky Way contains $\\sim2\\times10^9\\,$M$_\\odot$ of gas, so the CMZ holds roughly a few percent of this. The WMAP analysis shows the CMZ also contains a few percent of the ionising photons, and hence star formation, in the Galaxy. Both the volumetric and surface density SF relations predict that a given mass of gas will form stars more rapidly if the respective densities are larger. Yet the gas in the CMZ, which has a much higher surface and volume density than any comparable mass of gas in the disk of the Milky Way, forms stars at a rate proportional to the ratio of gas in the CMZ to that in the Milky Way.  Something is required to slow the rate of SF in the CMZ compared to that in the rest of the Milky Way. An additional support mechanism, not taken into account in either the column density threshold or volumetric SF relations, may be responsible for inhibiting the SF. One potential solution is therefore an additional term or threshold in the proposed SF relations. The most noticeable difference between clouds in the CMZ and the rest of the Milky Way, apart from the 1 to 2 order of magnitude larger volume density, is the order of magnitude larger internal cloud velocity dispersion \\citep{morris_serabyn1996,ferriere2007}. We thus conclude it is likely that the relevant term is related to the additional turbulent energy in the gas providing support against gravitational collapse. A simple way to parametrise this is through the linewidth ratio, $\\Delta V_{\\rm ratio} = \\Delta V/\\Delta V_0$, where $\\Delta V_0$ is the typical internal cloud velocity dispersion in disk molecular clouds. The Schmidt law could then be re-expressed either in terms of surface density, $\\Sigma_{SFR} \\propto (\\Sigma_{\\rm gas})^\\alpha/(\\Delta V_{\\rm ratio})^b$ [$\\alpha \\sim$ 1 to 1.4; $b\\sim1$], or mass of gas above the density threshold, $M_{\\rm dense}$, through, $\\Sigma_{SFR} \\propto M_{\\rm dense} /(\\Delta V_{\\rm ratio})^b$. This would reconcile the star formation in the CMZ with that in the rest of the Galaxy. Theorists have wrestled with this before \\citep[e.g.][]{krumholz_mckee2005,padoan_nordlund2011,kdm2012} and we are currently seeking to test these scenarios.  However, even if the extreme environmental conditions in the CMZ do inhibit SF, they do not stop it entirely. The CMZ contains Sgr B2, one of the most extreme cluster forming regions in the Galaxy. It also contains high-mass star clusters like the Arches and Quintuplet, and at least one molecular cloud which appears to be the progenitor of such massive stellar clusters \\citep{longmore2012_brick,bressert2012b}. Evidence exists of episodic SF events in the CMZ \\citep{sofue_handa1984,yusef-zadeh2009,su2010,bland-hawthorn_cohen2003} and mechanisms exist to explain how such episodic SF can occur. Gas in barred spiral galaxies like the Milky Way is funnelled from the disk through the bar to the GC \\citep{kormendy_kennicutt2004,sheth2005}. If gas is continually fed from the disk to the GC, and the environmental conditions impose a higher threshold for SF to occur, the gas might build up until reaching a critical point before undergoing a burst of SF.  Although clouds near the Galactic centre in the Milky Way and other galaxies may only represent a small fractional volume of a galaxy, they can contribute a significant fraction of the total dense molecular gas. In terms of dense gas mass and environmental conditions, the Galactic centre also acts as a bridge between local SF regions in our Galaxy and SF environments in external galaxies. Understanding why such large reservoirs of dense gas deviate from commonly assumed SF relations is of fundamental importance and may help in the quest to understand SF in more extreme (dense) environments, like those found in interacting galaxies and at earlier epochs of the Universe."
    },
    "1208/1208.6065_arXiv.txt": {
        "abstract": "For the past forty years the search for dark matter has been one of the primary foci of astrophysics, although there has yet to be any direct evidence for its existence \\citep{Porter2011}. Indirect evidence for the existence of dark matter is largely rooted in the rotational speeds of stars within their host galaxies, where, instead of having a $\\sim r^{-1/2}$ radial dependence, stars appear to have orbital speeds independent of their distance from the galactic center, which led to proposed existence of dark matter \\citep{Porter2011, Peebles1993}. We propose an alternate explanation for the observed stellar motions within galaxies, combining the standard treatment of a fluid-like spacetime with the possibility of a ``bulk flow'' of mass through the Universe. The differential ``flow'' of spacetime could generate vorticies capable of providing the ``perceived'' rotational speeds in excess of those predicted by Newtonian mechanics. Although a more detailed analysis of our theory is forthcoming, we find a crude ``order of magnitude'' calculation can explain this phenomena. We also find that this can be used to explain the graviational lensing observed around globular clusters like ``Bullet Cluster''. ",
        "introduction": "\\label{sec:Intro}  In the pursuit of determining a model that accurately predicts the past, present, and future of the evolution of the Universe, physicists have generated a range of possible candidates. Currently, the most generally accepted being the $\\Lambda$CDM model, which contains, among others, parameters dealing with the existence of a cosmological constant($\\Lambda$) and cold dark matter (CDM). Furthermore, the existence of the dark matter component of this cosmology, and others, is not that which is generally considered contentious. Dark matter has rather become somewhat of a staple in the diet of cosmologies. However, there are observational reasons to give pause to the assumed existence of Universal cold dark matter, which then should lead us to question whether or not there are other alternative models.  Models of dark matter succeed in accounting for the galactic rotation curves observed throughout the Universe, by increasing the mass of the galaxy beyond the observed. There are, however, simple problems with the dark matter halo model that have yet to be fully explained (e.g. \\citep{Klypin1999, Garbari2011, Karachentsev2012, Poitras2012, Sluse2012}). One of these problems is the disparity between the observed (stellar) mass function, usually defined in terms of the Schechter function, and the theoretically expected cosmological halo mass function \\citep{Moster2010}. One of the defining problems of galactic formation and evolution is determining the origins of this disparity. This is an example of how our understanding of dark matter (or lack there of) is still grounds for much debate. However, if it may be possible to utilise a different model for the origin of galactic formation then it is possible that some of these questions may be answered.  Recently, there has been some observational evidence for the ``bulk flow'' of matter through the Universe \\citep{Benson2003, Bhattacharya2008, Feldman2010, Osborne2011, Turnbull2012}. If these measurements prove to be true, then the nature of this flow is of interest beyond that of the distribution of matter in the Universe. Specifically relevant to this discussion is the interpretation that the ``flow'' observed is caused not by an {\\em en masse} transit of matter through the Universe, but rather by the motion of spacetime itself. Whilst this concept is indeed foreign it can be considered somewhat preferential to the former case, from an isotropic viewpoint, as the motions of objects in the Universe need not be preferentially oriented in this regime. More importantly, variation in the ``bulk flow'' of spacetime fluid through the Universe could produce eddies in spacetime and provide the additional unexplained velocity to rotational speeds of stars beyond the central bulge of galaxies.  As the intention of this paper is to merely propose an alternative theoretical explanation for observations consistent with the existence of dark matter the structure is as follows: \\S~\\ref{sec:spacetime} describes the treatment of spacetime as a fluid. Sections \\ref{sec:classical} and \\ref{sec:relativistic} discuss classical fluid dynamics and relativistic fluid dynamics and how votices in the spacetime ``fluid'' can produce observations consisten with dark matter. Finally we present concluding remarks and propose future work in \\S~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} If the observed effects attributed to dark matter are indeed caused by the turbulent flow of spacetime, then we can simply hypothesise that any galaxy in a cluster formed in this way should all rotate in the same direction. Albeit only significant to $1.6 \\sigma$, evidence for this effect has been detected by \\citet{Longo2011}. There is no reason to expect that this observation would also be caused in the dark matter Regime for reasons other than chance. As this is a simple method to determine if turbulent spacetime flows may be the cause of galactic rotational curves, Doppler measurements of stellar velocities in galaxies are extremely important. Furthermore, if rotational curves for distant galaxies can be found, isotropy measurements could serve as an additional constraint for the validity of this theory.  Finally, we find that chaotic flows could exist on both ``large'' and ``small'' scales. Large scale turbulence is dominated by the ``flow'' velocity and the uniformity thereof. In this context, ``large'' scale turbulence would be on a galactic scale, with ``large'' eddies being comparable to the size of a galaxy, which could be used to explain why galaxies are not ``sheared'' apart. Small scale turbulence is dominated by the viscosity of the flow, which would most likely be caused by the gravitational attraction of masses present in the region of the eddy, as mentioned previously in relation to the galactic rotation curves. The scale of a ``small eddy'' would be comparable to that of stellar clusters. Since these clusters are still affected by distortions in spacetime, this could explain the observed gravitational lensing caused by some stellar clusters that could not be explained by modified Newtonian dynamics (MOND) \\citep{Ibata2011}.  This concept can, therefore, effectively explain the major observations that lead to the introduction of dark matter, and removes the need for the existence of a massive dark matter halo about galaxies. If this theory is correct, galaxies co-located within a differentially moving frame, should all rotate in the same direction (i.e. same chirality). Additionally, when observed from an external reference frame, there should be variations in the local spacetime of a galaxy $\\delta g u_t \\sim 10^{-4} \\km \\s^{-1} \\Lyr^{-1}$ where $u_t$ is the observed ``flow'' of the surrounding galaxies.  Finally can now draw some different, interesting, conclusions from the disparity between the observed stellar mass function and the cosmological halo mass function. Using the model we propose here, the origins of galactic evolution lie not in vast halos of dark matter, but rather in the turbulence of spacetime. The turbulence itself traces back its origins to spatial and temporal variations in the motion of spacetime. If we are able to conceptualise this method of seeding galaxies, then we can also recognise that the formation of the vortices in which galaxies originate is highly dependent on the mechanics of the local spacetime. This, when interpreted simply, means that the expected distribution of galactic mass (i.e. the observed stellar mass function) should not be conformal to a simplistic power law (i.e. the cosmological halo mass function), but rather should be more complex in nature. The true distribution of vortex sizes in a field of uniform variation should be expected to be inately coupled with galactic mass. This may well explain the shape of the Schechter function, and why galactic mass function does not follow a simple power law, without the need for complex models to explain the reduction of star formation in the low and high mass dark matter halo regimes.  As stated above, many of these conclusion are too complicated to present in any detail within this paper, of which will be presented in following papers. The purpose of this paper is to merely propose an alternate explanation for the observed effects of dark matter and a possible explanation for why dark matter is not consistent with some other theories and observation"
    },
    "1208/1208.4851_arXiv.txt": {
        "abstract": "Size differences of $\\approx 20\\%$ between red (metal-rich) and blue (metal-poor) sub-populations of globular clusters have been observed, generating an ongoing debate as to weather these originate from projection effects or the difference in metallicity. We present direct $N$-body simulations of metal-rich and metal-poor stellar populations evolved to study the effects of metallicity on cluster evolution. The models start with $N = 100\\,000$ stars and include primordial binaries. We also take metallicity dependent stellar evolution and an external tidal field into account. We find no significant difference for the half-mass radii of those models, indicating that the clusters are structurally similar. However, utilizing observational tools to fit half-\\emph{light} (or effective) radii confirms that metallicity effects related to stellar evolution combined with dynamical effects such as mass segregation produce an apparent size difference of $17\\%$ on average. The metallicity effect on the overall cluster luminosity also leads to higher mass-to-light ratios for metal-rich clusters. ",
        "introduction": "Globular clusters (GCs) are substantial components of galaxies and found in populations of up to thousands in giant elliptical galaxies \\citep{Peng2011}. The Milky-Way (MW) hosts a GC population of $157$ confirmed clusters (\\citealt{Harris1996}, 2010 edition), with new clusters still being discovered (e.g. \\citealt{Minniti2011}). These clusters live within the bulge as well as the halo of the Galaxy and can - in contrast to star clusters beyond the Local Group - easily be resolved in ground-based observations. In general, the GC systems of galaxies tend to appear in two sub-populations: a blue and a red component \\citep{Zinn1985}. Although the metallicity cannot be inferred directly from the cluster colour due to the age-metallicity degeneracy \\citep{Worthey1994}, it has been well accepted that the blue clusters are metal-poor, whereas the red ones are metal-rich. Both sub-populations are old (e.g. \\citealt{MarinFranch2009}), with a trend for the red clusters to be more centrally concentrated within their host galaxy's potential than their blue counterparts (\\citealt{Kinman1959, BS2006}, also see Fig. \\ref{fig:fig1}).  The ability of the Hubble Space Telescope to partially resolve globular clusters even beyond the Local Group has lead to the finding that i) GCs have mean half-light radii $r_{\\rm{hl}}=3\\,$pc \\citep{Jordan2005} and ii) red clusters are on average $\\approx 17-30\\%$ smaller than their metal poor counterparts \\citep{Kundu1998, Jordan2005, Woodley2010}. Several explanations for this phenomena have been proposed: projection effects and the influence of the tidal field \\citep{Larsen2001} or a combined effect of mass segregation and the dependence of main-sequence lifetimes on metallicity \\citep{Jordan2004, Jordan2005}. Whether either of those effects are dominating or a combination of both can only be investigated through direct star cluster simulations where three-dimensional galactocentric distances are known and stellar evolution is included in the dynamical evolution of the cluster.  The effects of metallicity on the evolution of a single star manifests itself as a different rate of stellar evolution, which is accompanied by a different mass-loss rate and hence ultimately affects the stars lifetime and remnant mass (see Section \\ref{Zeff}). In general, low metallicity stars evolve \\emph{faster} along the main sequence than their high metallicity counterparts \\citep{SSE}. For a bound system such as a star cluster, the increased mass-loss rate can lead to a lower cluster mass and hence a lower escape velocity and. This in turn can produce a stronger increase in radius for the metal-poor cluster. At later stages, this might also lead to postponed core-collapse for the low metallicity cluster. Both effects could lead to a larger measured cluster size. A preliminary study along these lines has been carried out by \\citet{Hurley2004} for open clusters. They showed that an increased escape rate for the metal-rich clusters owing to earlier core-collapse acts to cancel these effects resulting in only a $10 \\%$ difference in cluster lifetime for metal-poor versus metal-rich cases - within the statistical noise of fluctuating results from one simulation to another. However, several aspects of our new simulations differ from this preliminary study. Among those are an adjusted binary fraction for GCs and an improved tidal field. Most importantly we also use a higher initial number of stars $N_{\\rm{i}}$, bringing the $N$-body models into the GC regime. This ultimately leads to an increase in cluster lifetime and hence not necessarily core-collapse or depletion of stars within a Hubble time.  In this work, we make use of a set of star cluster simulations evolved with the direct $N$-body code \\texttt{NBODY6} \\citep{Aarseth1999, Aarseth2003} to study the effects of metallicity on star cluster dynamics, evolution and size (i.e. effective radius) to answer the question if metallicity alone could reproduce the observed size difference. We measure the sizes of these clusters along their evolutionary track with methods used both in observations and theory.  Recently \\citet{Downing2012} has published a set of Monte-Carlo models exploring the origin of the observed size difference between metal-rich and metal-poor GCs, which provides an excellent comparison for our work. This follows on from the $N$-body models of \\citet{Schulman2012}, who investigated the evolution of half-mass radius with metallicity in small-$N$ clusters. Similarly to \\citet{Downing2012}, we shall be careful to make a distinction between the actual size of a star cluster, represented by the half-mass radius (which we shall denote as $r_{\\rm{50\\%}}$, i.e. the $50\\%$ Lagrangian radius), and the observationally determined size (the half-light or effective radius, $r_{\\rm{eff}}$).  This paper is structured as follows. We introduce the differences in stellar evolution depending on metallicity in the next Section. In Section \\ref{method}, we describe our simulation method and the models we have chosen to evolve. In Section \\ref{evolution}, we analyze the evolution of cluster mass, binary fraction, luminosity, half-light radius and mass-to-light ratio which is followed by discussion and conclusions. ",
        "conclusions": "We have measured the sizes of GC models with different metallicity, evolved with the direct $N$-body code \\texttt{NBODY6}. All clusters start their evolution with $105\\,000$ stars and a mass of $\\approx 6 \\times 10^4~M_{\\odot}$. We find no size differences with metallicity when measuring sizes by means of the half-mass radius or other mass-weighted radii, with the exception that lower remnant masses for high-$Z$ stars cause the $N$-body core radius to fluctuate less. This indicates, that there is no structural difference between clusters of low and high metallicity. Even though the mass-loss rates of low-$Z$ stars are higher, especially in the initial stages of evolution, a consequently lower escape velocity and higher average remnant mass cancels this effect, leading to no overall size difference. In accordance with this, we also find that the number of stars and cluster mass remaining at a particular time do not vary noticeably with the metallicity of the cluster.  \\citet{Schulman2012} evolved $N$-body models starting with $N = 8\\,192$ stars and different metallicities to find a size difference between metal-poor and metal-rich clusters, in terms of the half-mass radius. This is in disagreement with our results and those of the Monte Carlo models of \\citet{Downing2012}. The \\citet{Schulman2012} models were evolved with some softening so that the effects of close binaries were not included. They were evolved to a dynamical age of $5 \\, t_{\\rm rh}$ which translated to physical ages in the range of $100 - 500\\,$Myr for the small-$N$ models. The claim is that the results should be applicable to larger clusters, including GCs, because the impact of different stellar evolution and mass-loss histories at various $Z$ will not depend on $N$, and also because they performed models in the range of $1\\,024$ to $16\\,384$ stars that showed similar half-mass radius evolution. We would counter that as the MS lifetime of a MS turn-off star changes with age and the half-mass relaxation timescale of a cluster varies with $N$, it is not at all obvious that the interplay between stellar evolution and cluster dynamics will scale in a straightforward manner. Indeed, our models here and the open cluster models of \\citet{Hurley2004} with $N \\sim 30\\,000$, both show that the half-mass radius of metal-rich models can be smaller than that of the metal-poor models at early times (see Fig. \\ref{fig:fig6}) but that the difference is erased or even reversed later in the evolution. Factors including different core-collapse times, the stellar evolution of low-mass stars as a function of metallicity (particularly for globular clusters with ages of $10\\,$Gyr or more) and different remnant masses need to be taken into account to gain the full picture. Furthermore, statistical fluctuations are generally prevalent in small-$N$ simulations and it can be necessary to average the results of many instances to establish true behavior (e.g. \\citealt{Kuepper2008}). Our models presented here are at the lower edge of the GC mass function but even for these we would suggest that larger models again are desired before making any final judgment about the size measurements of GCs in general. However, our agreement with the large-scale Monte Carlo models of \\citet{Downing2012}, performed with $5 \\times 10^5$ stars, on the issue of half-mass radius variation (or non-variation) with metallicity is reassuring.  In contrast to the evolution of the half-mass radius, we find that the half-\\emph{light} (or effective) radius does vary with metallicity. We find that blue, metal-poor clusters can appear on average $17\\%$ larger than red, metal-rich clusters, with even larger differences possible when comparing individual models. This is in agreement with observations of extra-galactic GC systems, where size differences of $17-30\\%$ \\citep{Larsen2001, Jordan2005, Woodley2010} have been found. It is also in agreement with the Monte Carlo models of \\citet{Downing2012}. Indeed, our $N$-body models and these Monte Carlo models provide excellent independent validation of the main result -- that the observed size differences in GCs are likely caused by the interplay of stellar evolution and mass segregation. Stellar evolution causes low-$Z$ stars to be brighter than their high-$Z$ counterparts while mass segregation causes the most massive remnants to sink to the centre. Successively more massive remnants in low-$Z$ clusters leads to a steeper surface brightness profile for high-$Z$ clusters. The overall mass segregation is similar for metal-poor and metal-rich clusters but more effective in the luminous stars for high-$Z$ clusters owing to a higher main-sequence turnoff mass. This is in excellent agreement with the predictions of \\citet{Jordan2004} using multi-mass Mitchie-King models to estimate the size difference between blue and red GCs, finding a difference of $14\\%$ due to the combined effect of mass-segregation and stellar evolution.  The apparent size difference does have a dependence on the treatment of remnants. When ejecting all NSs and BHs, no significant size difference (half-light radius) is found, partly owing to the fact that one of the variations with metallicity (remnant masses) has been negated. When we retain $\\approx 5\\%$ of the NSs and BHs arising from the primordial population, our results are in general agreement with the \\citet{Downing2012} models that retained large numbers of BHs. While there are uncertainties in the retention fractions for NSs and BHs, there are also uncertainties for the masses of remnant BHs. We have used the stellar evolution wing mass-loss prescriptions from \\citet{SSE}, while improved, $Z$ dependent mass-loss rates are now available \\citep{Vink2001}. However, the resulting differences for BH masses are most apparent for stars above $40\\,M_{\\odot}$ \\citep{Belczynski2010}, while just a few stars are drawn from this mass range in the models presented here.  The average size difference of $17\\%$ implies that blue GCs do indeed appear larger as a result of metallicity effects. Since this is at the lower end of what is found in observations, other causes (such as projection effects) can also be expected to play a role. In the future we plan to extend our study by performing additional $N$-body simulations that explore parameters such as larger $N$, smaller initial size and differing initial density profiles, as well as different cluster orbits, to further understand the effects of cluster evolution and environment on measured sizes. Our spread of individual measurements in Fig. \\ref{fig:fig8} can be compared to extragalactic studies of GC systems as well as in the Milky Way, in which half-light radii of GCs are found to be distributed between $1$ to $8\\,$pc (e.g. \\citealt{Larsen2003} Fig. $4$, \\citealt{Spitler2006} Fig. $19$, \\citealt{Madrid2009} Fig. $10$). Since clusters of different masses and at different galactocentric distances are included in the observational samples, a larger scatter is expected than for our models (which currently give values between $2-6\\,$pc). We would expect the model spread to increase when we extend our study to include a range of cluster parameters.  In addition to the half-light radius, we have also analyzed the evolution of the mass-to-light ratio. When comparing cluster models evolved purely through stellar (but no dynamical) evolution with the thorough $N$-body models, there is little change in $M/L$. As seen before in \\citet{Baumgardt2003}, we find that $M/L$ increases with time, where dynamical interactions lead to a decrease in $M/L$ as low-mass stars (carrying a high mass-to-light ratio) are preferentially lost from the cluster. The decrease in overall cluster luminosity with time results in an increase of the mass-to-light ratio."
    },
    "1208/1208.4060_arXiv.txt": {
        "abstract": "{ We study the globular cluster (GC) systems in three representative fossil group galaxies: the nearest (NGC\\,6482), the prototype (NGC\\,1132) and the most massive known to date (ESO\\,306-017).  This is the first systematic study of GC systems in fossil groups. Using data obtained with the \\textit{Hubble Space Telescope} Advanced Camera for Surveys in the F475W and F850LP filters, we determine the GC color and magnitude distributions, surface number density profiles, and specific frequencies. In all three systems, the GC color distribution is bimodal, the GCs are spatially more extended than the starlight, and the red population is more concentrated than the blue. The specific frequencies seem to scale with the optical luminosities of the central galaxy and span a range similar to that of the normal bright elliptical galaxies in rich environments. We also analyze the galaxy surface brightness distributions to look for deviations from the best-fit S\\'ersic profiles; we find evidence of recent dynamical interaction in all three fossil group galaxies. Using X-ray data from the literature, we find that luminosity and metallicity appear to correlate with the number of GCs and their mean color, respectively. Interestingly, although NGC\\,6482 has the lowest mass and luminosity in our sample, its GC system has the reddest mean color, and the surrounding X-ray gas has the highest metallicity. }  {} ",
        "introduction": "The most accepted mechanism of galaxy formation is hierarchical assembly, whereby large structures are formed by the merging of smaller systems  (Press \\& Schechter 1974; De Lucia et al. 2006); nevertheless, there are some open questions and this scenario is debated (Nair et al. 2011, see also Sales et al. 2012). The hierarchical assembly is supported by the frequent observation of galactic interactions, and predicts that all the galaxies within a galaxy group or cluster will eventually merge into a single massive elliptical galaxy; this would be the last step of galaxy formation. Indeed, Ponman et al.\\ (1994) found an extreme system consisting of a giant elliptical galaxy surrounded by dwarf companions, all immersed in an extended X-ray halo. This system was interpreted as the end product of the merger of galaxies within a group, and was thus called a fossil group (FG). A formal definition was introduced by Jones et al.\\ (2003), according to which FGs are systems with $L_x\\ge10^{42}h_{50}^{2}\\rm{erg~s}^{-1}$, and an optical counterpart where the difference in magnitude between the first and the second brightest galaxies is ${\\Delta}m_R>2~$mag. Because of their regular X-ray morphologies and lack of obvious recent merger activity, FG galaxies are usually considered to be ancient and unperturbed systems (Khosroshahi et al. 2007). This is supported by numerical simulations that suggest FGs formed at early epochs ($z>1$) and afterwards evolved fairly quiescently (D'Onghia et al.\\ 2005; Dariush et al.\\ 2007; D{\\'{\\i}}az-Gim{\\'e}nez et al.\\ 2011).  At present, there is no consensus on the mass-to-light ($M/L$) ratios of FGs; measurements span from $\\sim100M_\\odot/L_\\odot$ (Khosroshahi et al.\\ 2004) to $\\sim1000M_\\odot/L_\\odot$ (Vikhlinin et al.\\ 1999; Yoshioka et al. 2004; Cypriano et al.\\ 2006) in the $R$-band, but there is a tendency towards higher values (Proctor et al. 2011; Eigenthaler \\& Zeilinger 2012). Such high $M/L$ values, together with the fact that FGs show unusually high $L_X / L_{opt}$ ratios compared with typical galaxy groups, appear inconsistent with the assumption that FGs are formed by the merging of the members in ordinary galaxy groups (see Fig.\\,2 in Khosroshahi et al.\\ 2007). Compact galaxy groups in the nearby universe tend to have lower masses and X-ray luminosities than observed in FGs, although there is evidence that high-mass compact groups may have been more common in the past and could represent the progenitors of today's FGs (Mendes de Oliveira \\& Carrasco 2007). In addition, the optical luminosities of the dominant giant ellipticals in FGs (Khosroshahi et al. 2006; Tovmassian 2010) are similar to those of brightest cluster galaxies (BCGs). For these reasons, it has been suggested that FGs are more similar to galaxy clusters, but simply lack other early-type galaxies with luminosities comparable to the central galaxy (Mendes de Oliveira et al.\\ 2009). It remains an open question whether FGs represent the final phase of most galaxy groups, or if they constitute a distinct class of objects which formed with an anomalous top-heavy luminosity function (Cypriano et al.\\ 2006; Cui et al.\\ 2011; see also M{\\'e}ndez-Abreu et al.\\ 2012).  Globular clusters (GC) systems are a powerful tool to study galaxy assembly (Fall \\& Rees 1985; Harris 1991; Forbes et al.\\ 1997; West et al. 2004; Brodie \\& Strader 2006). GCs are among the oldest objects in the universe, and dense enough to survive galactic interactions. Most of them have ages older than 10\\,Gyr (Cohen et al. 1998; Puzia et al. 2005), and therefore probably formed before or during galaxy assembly.  The correlation of the optical luminosity of a galaxy with the mean metallicities of the field stars and GCs (van den Bergh 1975; Brodie \\& Huchra 1991; Forbes et al.\\ 1996; C\\^ot\\'e et al.\\ 1998) supports the idea that GC and galaxy formation are strongly linked.  Nevertheless, there is a shift in metallicity, in the sense that field stars are on average more metal-rich than GCs. Interestingly, Cohen et al. (1998) noted that the metallicity of the X-ray gas in M87 is similar to the metallicity of the field stars.  Thus, the GC population could constitute a record of the initial chemical enrichment of the parent galaxy.  The GC systems of massive early-type galaxies are generally bimodal in their optical color distributions (Zepf \\& Ashman 1993; Geisler et al.\\ 1996; Gebhardt \\& Kissler-Patig 1999). The two GC color components are commonly referred to as the metal-poor (blue) and metal-rich (red) subpopulations, and could indicate distinct major episodes of star formation. In the Milky Way, spectroscopic studies reveal two populations of GCs with distinct metallicities (Zinn 1985).  However, the situation is less clear for giant ellipticals and other massive galaxies with more complex formation histories, which may simply have broad metallicity distributions with a nonlinear dependence of color on metallicity (Yoon et al.\\ 2006, 2011; Blakeslee et al.\\ 2012; Chies-Santos et al. 2012). Another important property of GC systems is their luminosity function (GCLF). The GCLF is well described by a Gaussian with peak at ${M_V}^0\\approx-7.4$, and the dispersion for giant ellipticals with well-populated GC systems is ${\\sigma_V}^{\\rm GCLF}\\approx1.4$ (Harris 1991). The general homogeneity of the GCLF is unexpected {\\it a priori}, given that GC destruction mechanisms should depend on the environment (McLaughlin \\& Fall 2008). However, recent studies show that the GCLF dispersion depends on galaxy luminosity (\\jordan\\ et al.\\ 2007b; Villegas et al.\\ 2010).  The most basic measure of a GC system is its richness, or the total number of GCs in the system. Harris \\& van~den~Bergh (1981) introduced the concept of {\\it specific frequency}, ${S_N}$, as the number of GCs per unit galaxy luminosity, normalized to a galaxy with absolute $V$ magnitude $M_V = -15$. They reported values in the range 2$< S_N<$10 for elliptical galaxies. However, they also already noted that M\\,87, the central galaxy in the Virgo cluster, had an outstandingly large value of $S_N\\sim20$ (a more recent value from Peng et al. 2008 is $S_N=12.6\\pm0.8$); hence, they suggested that environment may play a role in GC formation. Later, other high-$S_N$ galaxies were identified, most of them being either BCGs or second brightest galaxies in clusters; often, these galaxies were classified as type cD, however not all galaxies classified as cD have high $S_N$ (Jord\\'an et~al.\\ 2004). It is interesting to note that, although BCGs have quite a uniform luminosity, to the point of being considered standard candles (Postman \\& Lauer 1995), they have a wide range in $S_N$. Conversely, Blakeslee et al.\\ (1997) found a correlation between $S_N$ of the BCG and the velocity dispersion and X-ray luminosity of the cluster.  They therefore suggested that the number of GCs scales roughly with cluster mass, a fact that suggests that galaxies with high values of $S_N$, like M87, are not anomalously rich in GCs but, rather, underluminous as a consequence of the lower overall star formation efficiency in more massive and denser systems (Blakeslee 1999; Peng et al.\\ 2008).  Hence, given that FGs appear to be ancient, highly luminous, but relatively unperturbed systems with origins that remain poorly understood (identified less than two decades ago), we study for the first time the GC populations of the dominant elliptical.  We also compare our measured GC properties with X-ray properties obtained from the literature, and look for photometric irregularities that could be signs of recent galactic interactions. In order to have a representative sample we choose: the nearest, NGC\\,6482 ($z=0.013$), the prototype, NGC\\,1132 ($z=0.023$), and the most massive known to date, ESO\\,306-017 ($z=0.036$), shown in Fig.\\,\\ref{pretty_pictures}. In the following section, we describe the observations, data processing, and GC selection criteria; in Sec.~3, we analyze the photometry and present our measurements; in Sec.~4, we discuss the implications of our results in the larger context of FG and early-type galaxy formation scenarios. The final section summarizes our conclusions. Throughout, we estimate distances assuming $H_0=71$ km~s$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{section_conclusions}  The main results of this study are:  \\begin{enumerate} \\item{We detected 369, 1410 and 1918 GCs down to magnitude $z$=24.7, 26.2 and 26.2 in NGC\\,6482, NGC\\,1132 and ESO\\,306-017, respectively; after completeness correction and assuming a Gaussian GCLF with expected ${M^0}_g=-7.2$ and $\\sigma_{GCLF}=1.4$, the number of GCs in each FOV are: 1140$\\pm$87, 3613$\\pm$295 and 11577$\\pm$1046, respectively.\\\\ } \\item{The GC color distributions for all three FGs are better described by a bimodal, rather than unimodal  Gaussian model. For the heteroscedastic cases, NGC\\,1132 and ESO\\,306-017, the dispersion of the red population is wider than the blue. The mean $g{-}z$ colors are: 1.26, 1.11 and 1.20 for NGC\\,6482, NGC\\,1132 and ESO\\,306-017, respectively.  The mean color for NGC\\,6482 is unusually red; it is interesting that this is also the FG with the highest X-ray gas metallicity.\\\\ } \\item{For all three FGs, the spatial distribution of the starlight is more concentrated than that of the GCs, and the blue GCs are more extended than the red ones, similar to the case for other ellipticals.\\\\ } \\item{The derived values of $S_N$ are: $1.9\\pm0.1$, $3.1\\pm0.3$ and $6.3\\pm0.6$ for NGC\\,6482, NGC\\,1132 and ESO\\,306-017, respectively. These span the full range for normal ellipticals in the Virgo cluster.  The results are consistent with both the total number of GCs and $S_N$ increasing with the optical luminosity of the galaxy and the X-ray luminosity from the intra-group gas. \\\\ } \\item{ From analysis of the surface brightness distributions, we found evidence of recent interactions, particularly in ESO\\,306-017, which shows a tidal feature coincident with the {\\it finger} reported previously in X-rays, and NGC\\,1132, which has shell-like structure. While NGC\\,6482 is well-described by a standard \\sersic\\ profile, the two brighter galaxies are better fitted by core-\\sersic\\ models. All three galaxies contain central dust features. These observations are consistent with numerical simulations indicating that signs of recent merging should be fairly common in first-ranked FG galaxies (D\\'iaz-Gim\\'enez et al.\\ 2008), and they suggest that the paradigm of FGs as relaxed, undisturbed systems needs to be reconsidered.  \\\\ } \\end{enumerate}   Larger samples of GC systems in FGs are needed in order to make more definite conclusions. However, overall we conclude that the GC properties (colors, spatial distributions, specific frequencies) in FG central galaxies are generally similar to those seen in other giant ellipticals, which mainly reside in clusters.  Although the environments differ, this study suggests that the GC systems formed under very similar conditions. These results might therefore be taken as a confirmation that the same basic formation processes are responsible for the buildup of massive early-type galaxies in all environments."
    },
    "1208/1208.3462_arXiv.txt": {
        "abstract": "Space-based coronagraphs for future earth-like planet detection will require focal plane wavefront control techniques to achieve the necessary contrast levels. These correction algorithms are iterative and the control methods require an estimate of the electric field at the science camera, which requires nearly all of the images taken for the correction. We demonstrate a Kalman filter estimator that uses prior knowledge to create the estimate of the electric field, dramatically reducing the number of exposures required to estimate the image plane electric field. In addition to a significant reduction in exposures, we discuss the relative merit of this algorithm to other estimation schemes, particularly in regard to estimate error and covariance. As part of the reduction in exposures we also discuss a novel approach to generating the diversity required for estimating the field in the image plane. This uses the stroke minimization control algorithm to choose the probe shapes on the deformable mirrors, adding a degree of optimality to the problem and once again reducing the total number of exposures required for correction. Choosing probe shapes has been largely unexplored up to this point and is critical to producing a well posed set of measurements for the estimate. Ultimately the filter will lead to an adaptive algorithm which can estimate physical parameters in the laboratory and optimize estimation. ",
        "introduction": "\\label{sec:intro} The desire to directly image extrasolar terrestrial planets has motivated much research into space-based missions. One approach proposed for direct imaging in visible to near-infrared light is a coronagraph, which use internal masks and stops to change the point spread function of the telescope, creating regions in the image of high contrast where a dim planet can be seen.  Coronagraphs possess an extreme sensitivity  to wavefront aberrations generated by the errors in the system optics (occulters are immune to this problem because the starlight never enters the telescope).  This necessitates wavefront control algorithms to correct for the aberrations and relax manufacturing tolerances and stability requirements within the observatory. In this paper we discuss the challenges associated with wavefront estimation and control in a coronagraphic imager. Advances in these correction algorithms have primarily been focused on development of the controller, by choosing some criterion that decides how best to suppress aberrations given an estimate of the electric field at that point in time. These estimators do not utilize prior knowledge of the electric field estimate, and as such require a large number of images to reconstruct the estimate. By utilizing prior estimates and the control history we develop a method that requires fewer images to update the estimate, thus improving the efficiency of the correction algorithm. ",
        "conclusions": "In this paper we have demonstrated a discrete time extended Kalman filter to estimate the image plane electric field in closed loop. This type of progress is critical for improving the efficiency of future coronagraphic missions, thereby maximizing the likelihood of planetary detections. We demonstrate the fastest suppression to date in the Princeton HCIL by only requiring a single measurement at each iteration, currently requiring $\\approx$30\\% of the original set of images for estimation. Faster algorithms also makes focal plane estimation techniques more feasible for ground-based coronagraphic instruments. The closed loop nature of the estimator also provides a more stable to measurement because a measurement update with poor signal-to-noise does not adversely affect the covariance of the state estimate. We have also shown with the single measurement update that not all probe shapes are best for estimation, motivating us to try using the control shape as the probing function. Preliminary results for using the control shape as a probe signal is promising, and may further reduce the number of required exposures to one image per iteration. The Kalman filter also opens up the possibility of adaptive control techniques to learn laboratory physical parameters and bias estimation to gain certainty in planetary detection using only the control history."
    },
    "1208/1208.3979_arXiv.txt": {
        "abstract": "{As a part of the \\emph{Herschel} key programme PRISMAS, we have used the \\emph{Herschel}-HIFI instrument    to observe interstellar nitrogen hydrides along the sight-lines  towards eight  high-mass star-forming regions   in order to elucidate the production pathways leading to nitrogen-bearing species in diffuse gas. Here, we report  observations towards W49N of the NH \\mbox{$N$\\,=\\,1\\,--\\,0},  \\mbox{$J$\\,=\\,2\\,--\\,1}, and  \\mbox{$J$\\,=\\,1\\,--\\,0}, \\mbox{ortho-NH$_2$} \\mbox{$N_{K_a, K_c}J$\\,=\\,1$_{1,1}3/2$\\,--\\,0$_{0,0}$1/2}, \\mbox{ortho-NH$_3$} \\mbox{$J_K$\\,=\\,1$_0$\\,--\\,0$_0$} and \\mbox{2$_0$\\,--\\,1$_0$},  \\mbox{para-NH$_3$} \\mbox{$J_{K}$\\,=\\,2$_{1}$\\,--\\,1$_{1}$} transitions, and unsuccessful searches for NH$^+$. All detections show absorption   by foreground material over a wide range of velocities, as well as absorption associated directly with the hot-core source itself. As in the previously published observations towards G10.6$-$0.4, the NH,  NH$_2$ and NH$_3$  spectra towards W49N show strikingly similar and non-saturated absorption features. We decompose the absorption of the foreground material towards W49N  into different velocity components in order to investigate whether the relative abundances vary among the velocity components, and, in addition,  we   re-analyse the   absorption lines towards G10.6$-$0.4~in the same manner. Abundances, with respect to molecular hydrogen,  in each velocity component are estimated using CH, which is found to correlate with H$_2$ in the solar neighbourhood diffuse gas.  \tThe analysis   points to a co-existence of the nitrogen hydrides in   diffuse or translucent  interstellar gas with a high molecular fraction. Towards both sources, we find   that NH is  always at least as abundant as both \\mbox{o-NH$_2$} and  \\mbox{o-NH$_3$}, in  sharp  contrast to previous results for dark clouds. We find relatively constant \\mbox{$N$(NH)/$N$(o-NH$_3$)} and \\mbox{$N$(o-NH$_2$)/$N$(o-NH$_3$)} ratios with mean values   of 3.2 and 1.9 towards W49N, and 5.4 and 2.2 towards G10.6$-$0.4, respectively. The    mean abundance of o-NH$_3$ is  $\\sim$2$\\times$10$^{-9}$  towards both sources. The nitrogen hydrides also show  linear correlations with    CN and HNC towards both sources, and looser correlations with CH. The upper limits on the  NH$^+$ abundance  indicate  column densities   \\mbox{$\\lesssim$2\\,--\\,14\\,\\%} of $N$(NH), which is in contrast to the behaviour of the abundances of CH$^+$ and OH$^+$ relative to the values determined for  the corresponding neutrals CH and OH. Surprisingly low values of the ammonia  ortho-to-para  ratio are found in both sources,  \\mbox{$\\approx$\\,0.5\\,--\\,0.7$\\pm$0.1}, in the strongest absorption components.  This result cannot be explained by current models as we had expected to find a value of unity or higher. } ",
        "introduction": "Nitrogen is among the six most abundant elements in the universe and, despite its fundamental role in the   chemistry of  molecules connected with life, the chemical network of nitrogen in the interstellar medium is still poorly understood due to severe difficulties to observe key molecules from the ground. Today  about 55 molecules containing nitrogen have been discovered in interstellar space and a few more in circumstellar envelopes\\footnote{www.astrochymist.org, www.astro.uni-koeln.de/cdms/molecules}. The major reservoir of nitrogen is believed to be in the form of atomic N and molecular N$_2$. The latter is extremely difficult to observe since it has no permanent dipole moment \\citep{2001ApJ...548..836S}. In dense molecular clouds N$_2$H$^+$ is instead often used as a tracer of  N$_2$, which remained undetected until   \\citet{2004Natur.429..636K} reported  far-ultraviolet observations towards HD\\,124314. The total abundance of nitrogen therefore still relies, to a high degree,   on chemical modelling and observations of nitrogen-bearing species other than  N$_2$.  In order to constrain the nitrogen formation pathways, observations of \\emph{nitrogen hydrides} are crucial since they are at the root of the nitrogen chemical network, appearing in its first steps in chains of reactions that lead to other more complex species. The  abundances of these species are thus key diagnostics for the nitrogen chemistry. The nitrogen hydrides are, however, also problematic to observe since their   ground state  rotational transitions   lie at sub-mm wavelengths and are thus very difficult, or impossible,  to reach from the ground. Key species, such as  imidogen (NH)  and amidogen (NH$_2$), have therefore previously not been widely observed,   and there is still no detection of the   NH$^+$ radical.    Although very few observations exist of NH and  NH$_2$~in interstellar space, they are well known in comets \\citep[e.g.][]{1941ApJ....94..320S, 1998Icar..136..268M, 1993ApJ...404..348F}, and have been observed  in stellar photospheres \\citep[e.g.][]{1969PASP...81..657S, 1989hra1.book.....F} via their electronic, vibration-rotation, and high rotational transitions. The first detection of interstellar NH was made by \\citet{1991ApJ...376L..49M} by optical absorption spectroscopy. Subsequent observations  by \\citet{1997MNRAS.291L..53C} and \\citet{2009MNRAS.400..392W} have yielded several lines of sight in  diffuse and translucent gas where column densities of NH, CH, CN, and H$_2$ have been directly measured. The average value of the column density ratio in these diffuse and translucent sight-lines is $N$(NH)/$N$(H$_2$)\\,=\\,3$\\times$10$^{-9}$.    Interstellar  NH$_2$ was first detected by \\citet{1993ApJ...416L..83V}  in absorption towards Sgr\\,B2 in three fine-structure components of the para-symmetry form of NH$_2$, the  \\mbox{$N_{K_a, K_c}$\\,=\\,$1_{1,0}-1_{0,1}$} transition, with partially resolved hyperfine structure at frequencies 461 to 469~GHz. The Infrared Space Observatory (ISO) was later used to observe   unresolved absorption lines of both ortho and para symmetry forms of  NH$_2$, as well as NH, towards this source  through the use of the long-wavelength spectrometer \\citep{2000ApJ...534L.199C, 2004ApJ...600..214G, 2007MNRAS.377.1122P}.    In contrast to NH and  NH$_2$, ammonia (NH$_3$) has been extensively observed for more than 40 years and was in fact the first polyatomic molecule to be identified in interstellar space \\citep{1968PhRvL..21.1701C} \t by means of its $J\\!=\\!K$ inversion transitions at cm wavelengths ($K$ is the quantum number of the projection of total angular momentum $J$ on the molecule's symmetry axis). The ammonia molecule has, however, two symmetries, like  NH$_2$,  which arise due to the possible orientations of the hydrogen spins and behave like two distinct species: \\mbox{ortho-NH$_3$} (all H spins parallel, $K$\\,=\\,3$n$ where $n$ is an integer $\\geq 0$) and \\mbox{para-NH$_3$} (not all H spins parallel, $K\\neq3n$). Important information about the ammonia formation pathways could   be inferred from observations of both symmetries. Unfortunately,  only the para form has relatively low-excitation transitions accessible from ground since the  $K$\\,=\\,0 ladder of energy levels has no inversion splitting and the $J_K\\!=\\!3_3$ inversion transitions' lower energy level is 122\\,K above ground. Ortho inversion lines ($K$\\,=\\,3$n$, $n \\ge 1$) can thus only be observed in relatively warm molecular gas. Only a few ammonia observations of  the cold diffuse interstellar gas exist, using para inversion lines \\citep{1990ApJS...72..303N,1994A&A...289..579T, 2006A&A...448..253L}   which leaves the ammonia formation mechanism poorly constrained in diffuse gas. The (0,0) ammonia ground state, with ortho symmetry, can  only be studied by \\mbox{sub-mm} rotational transitions. Observations of the fundamental rotational transition of \\mbox{ortho-NH$_3$} \\mbox{$J_K$\\,=\\,1$_0$\\,--\\,0$_0$} at 572~GHz, which has similar upper state energy as the inversion lines, but several orders of magnitudes higher critical densities,  thus  requires space telescopes. The Kuiper Airborne Observatory \\citep{1983ApJ...271L..27K} performed the first observations of the \\mbox{$J_K$\\,=\\,1$_0$\\,--\\,0$_0$} transition using heterodyne receivers, and later on the Odin satellite continued such observations in a number of environments,  for instance in photo-dissociation  and star-forming regions \\citep[e.g.][]{2003A&A...402L..69L,  2009A&A...494..637P}, circumstellar envelopes \\citep{2006ApJ...637..791H},  diffuse clouds in absorption towards Sgr\\,B2  \\citep{2010A&A...522A..19W},  and    in comets     \\citep{2007P&SS...55.1058B}.         \\begin{table*}[\\!htb] \\centering \\caption{Observed transitions towards W49N and G10.6$-$0.4. } \\begin{tabular} {lrccrrrllcc} \\hline\\hline \\noalign{\\smallskip} Species \t& Frequency\\tablefootmark{a}\t& Transition &\tBand\\tablefootmark{b} & $T_\\mathrm{sys}$\\tablefootmark{c}  & \\multicolumn{2}{c}{$t{_\\mathrm{int}}$\\tablefootmark{d}}   \t& \\multicolumn{2}{c} {$T_\\mathrm{C}$\\tablefootmark{e}} &\\multicolumn{2}{c}{$1\\sigma$/$T_\\mathrm{C}$\\tablefootmark{f}}    \\\\    \\noalign{\\smallskip} &&& && G10.6 & W49N & G10.6 & W49N  &  G10.6 & W49N   \\\\ & (GHz) &&  & (K) & (s)&  (s)& (K) & (K) &     \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip}  NH\\tablefootmark{g} &\t946.476\t& $N_J\\!=\\!1_0\\leftarrow0_1$ &\t  3b\t&  416&  234 &  116& 2.5 &3.4     &  0.021 &  0.027\\\\ &974.478\t&$N_J\\!=\\!1_2\\leftarrow0_1$      &\t4a & 338 \t &186   & 103\t&2.6& 3.9\t \t&  0.018  &  0.018   \\\\ o-NH$_2$\\tablefootmark{h}\t&\t952.578\t&  $N_{K_a,K_c} J = 1_{1,1} 3/2 \\gets 0_{0,0} 1/2$           &\t3b& 230  &92 &68& 2.6  &  3.6\t \t&0.018  &  0.017  \\\\  o-NH$_3$\\tablefootmark{i}\t&\t572.498&  $J_K$\\,=\\,1$_0 \\leftarrow 0_0$\t\t& 1b  & 87 & 1\\,024  &94& 0.61  &  0.93  \t\t \t&  0.013  &   0.025 \\\\   &\t1\\,214.859    & $J_{K}$\\,=\\,2$_{0} \\leftarrow 1_{0}$   & 5a  & 1\\,024   & 293 & 196 & 3.4 &  5.2  \t  & 0.032  &  0.025 \\\\ p-NH$_3$\t \t&\t1\\,215.246   &$J_{K}$\\,=\\, 2$_{1}  \\leftarrow 1_{1}$     &5a & 1\\,024     &293&  196&  3.4  & 5.2    \t &0.032  &  0.025 \\\\  NH$^{+}$ & \t1\\,012.540&   $^2\\Pi_{1/2}$ $N$\\,=\\,1$\\leftarrow$1  $J$\\,=\\,3/2\\,$\\leftarrow$\\,1/2 \t& 4a &  327\t &171& 91& 3.0 & 4.4    &0.013& 0.013  \\\\   \\noalign{\\smallskip} \\hline \\label{Table: transitions} \\end{tabular} \\tablefoot{ \\tablefoottext{a}{The frequencies refer  to the strongest hyperfine structure components.} \\tablefoottext{b}{HIFI consists of 7 different mixer bands and two  double sideband spectrometers. All transitions were observed in the upper sideband except NH$^+$.} \\tablefoottext{c}{System temperature.} \\tablefoottext{d}{The averaged on-source integration  time for each transition.} \\tablefoottext{e}{The single sideband (SSB) continuum intensity (the observed double sideband (DSB) continuum divided by two).} \\tablefoottext{f}{The rms noise, at a resolution of 1.1\\,MHz, divided by $T_\\mathrm{C}$.} \\tablefoottext{g}{The quantum numbers for the rotational transition $N_J$ are $F_1$\\,=\\,$I_H$\\,+\\,$J$ and $F$\\,=\\,$I_N$\\,+\\,$F_N$ \\citep{1997A&A...322L...1K}.} \\tablefoottext{h}{The quantum numbers for the rotational transition $N_{K_a, K_c}J$ are $F_1$\\,=\\,$I_N$\\,+\\,$J$ and $F$\\,=\\,$I_H$\\,+\\,$F_1$ \\citep{1999JMoSp.195..177M}.} \\tablefoottext{i}{The rotational energy is given by the two principal quantum numbers  ($J, K$), corresponding to the total angular momentum and its projection along the molecular axis \\citep[e.g.][]{1983ARA&A..21..239H, 1993JChPh..98.4662R}. } } \\end{table*}    With the launch  of the \\emph{Herschel} Space Observatory   \\citep{Pilbratt2010} in May 2009,    unique opportunities to perform observations of transitions between 157 and 625\\,$\\mu$m (0.48\\,--\\,1.9\\,THz)  became feasible  with the Heterodyne Instrument for the Far-Infrared \\citep[HIFI;][]{Graauw2010} owing to its very high sensitivity and spectral resolution. This allowed, for the first time, searches for spectrally resolved, rotational  transitions involving the ground states of  NH$^+$, NH,  NH$_2$, and NH$_3$~with the same instrument. Several  observations   have already  been reported. For instance, using  \\emph{Herschel}-HIFI, \\citet{2010A&A...521L..42B} found very high column densities of NH and ND in the cold envelope of the class\\,0  protostar IRAS16293-2422 (2$\\times$10$^{14}$\\,cm$^{-2}$~and $\\sim$1.3$\\times$10$^{14}$\\,cm$^{-2}$, respectively). \\citet{2010A&A...521L..52H} found  NH:NH$_2$:NH$_3$ abundance ratios  of $\\sim$5:1:300 towards the same source.    The PRISMAS\\footnote{http://astro.ens.fr/?PRISMAS}  key programme (PRobing InterStellar Molecules with Absorption line Studies) is targeting absorption lines in the line-of-sight towards eight bright sub-millimetre-wave continuum sources using \\emph{Herschel}-HIFI: G10.6$-$0.4~(W31C), W49N, W51, G34.3+0.1, DR21(OH), SgrA (+ 50~km~s$^{-1}$ cloud), G005.9-0.4 (W28A) and W33A.   High-resolution absorption line spectroscopy is generally a very sensitive  and model-independent method for measuring column densities of interstellar molecules, and a powerful tool to probe the diffuse interstellar gas clouds with no or little excitation.  The first results and analysis of absorption lines of nitrogen hydrides along the sight-line towards the massive star-forming region G10.6$-$0.4~(W31C) have already been presented in \\citet[][hereafter paper~I]{2010A&A...521L..45P}. Similar abundances with respect to the total amount of hydrogen $N_\\mathrm{H}$\\,=\\,2\\,$N(\\mathrm{H_2}$)+$N(\\mathrm{H}$),   were found for all three species:   approximately $6\\times 10^{-9}$,  $3\\times 10^{-9}$, and $3\\times 10^{-9}$ for NH, NH$_2$, and NH$_3$, respectively. They were estimated across the whole line-of-sight and using the high temperature ortho-to-para limits of three and one for NH$_2$ and NH$_3$, respectively.    NH$^+$ was not detected at a 1$\\sigma$ rms level of 74\\,mK with a resolution of 1.1~MHz. The abundance patterns that we see in diffuse molecular gas are thus clearly very different from those in IRAS16293-2422 and Sgr~B2, where NH:NH$_2$:NH$_3\\sim$1:10:100 and the fractional abundance of NH is a few times 10$^{-9}$. The Sgr~B2 results may, however, not be representative of cold dark clouds since  this source is very   complex and  atypical.    The unexpectedly high NH abundance has been difficult to explain with chemical models and both the NH and  NH$_2$ production by purely gas-phase processes  is inhibited by a lack of a sufficient source of N$^+$. The models fail to   simultaneously  predict the absolute and   relative abundances of the nitrogen hydrides. Typical steady state  \\emph{dark cloud} models ($n$\\,=\\,1$\\times$10$^{3}$\\,--\\,5$\\times$10$^{4}$~cm$^{-3}$, $T$\\,=\\,10\\,--\\,40~K, $A_\\mathrm{V}\\gtrsim10$),   predict an NH$_2$ abundance of \\mbox{(1-10)$\\times$10$^{-8}$}, an\t NH abundance 10 times lower, and an NH$_2$/NH$_3$~ratio of  \\mbox{0.3\\,--\\,1.5} for a wide range of assumptions \\citep[e.g.][]{1989ApJS...69..241L,1991A&AS...87..585M}, but these are not directly applicable to diffuse molecular gas. Examples of chemical models for diffuse cloud conditions are found in Fig.~A.1 and A.2 in paper~I. These models were also unable to explain the observed abundances and ratios. Processes on dust  grains have   previously been proposed as a way to increase the NH production \\citep{1991ApJ...376L..49M, 1993MNRAS.260..420W}. Such models, however, often predict up to 1\\,000 times more NH$_3$ than NH$_2$ \\citep[][paper~I]{1993MNRAS.263..589H}. The importance of grain surface chemistry in diffuse clouds is also not clear since water ice mantles have not been detected in diffuse gas, and  strong (UV) radiation fields counteract molecular formation on grains. Grain surface production of NH is therefore still debated, and would, if   true,   change our understanding of surface  chemistry in diffuse gas. On the other hand, if  grains indeed were unimportant in diffuse gas nitrogen chemistry it would imply that  either key gas-phase reactions must  have been overlooked,   or that the uncertainty of some reaction rates could make a difference. Both additional high-quality observations and chemical modelling are needed to solve this problem.  In this paper, we present new observations and  analyses of absorption lines in the line-of-sight towards the high-mass star-forming region  W49N, and, we also re-analyse the absorption towards G10.6$-$0.4~in more detail. W49N is one of the  most  luminous high-mass star-forming regions in the Galaxy ($\\sim$10$^7$\\,L$_\\odot$) with a core that contains more than a dozen ultra-compact \\ion{H}{II} regions \\citep{1984ApJ...283..632D, 1990ApJ...351..189D, 2000ApJ...540..308D}. It is located on the far side of the Galaxy at a distance of 11.4\\,kpc with Galactic coordinates $l$\\,=\\,43.17$^\\circ$ and $b$\\,=\\,0.012$^\\circ$, in one of the most massive giant molecular clouds ($\\sim$10$^6$\\,M$_\\odot$) in the Milky Way. The source velocity is about +8\\,km\\,s$^{-1}$~and the foreground gas  along the line-of-sight is detected at  \\mbox{$v_\\mathrm{LSR}\\approx 30-75$\\,km\\,s$^{-1}$}, revealing gas at two locations in the near and far side of the Sagittarius spiral arm around 40 and 60\\,km\\,s$^{-1}$~\\citep{1985ApJ...297..751D}. The  ultra-compact \\ion{H}{II} region G10.6$-$0.4~in the star-forming W31 complex  is an extremely luminous sub-millimetre and infrared continuum source. The source is located within the so-called 30\\,km\\,s$^{-1}$~arm at a kinematic distance of 4.8\\,kpc \\citep{2003ApJ...587..701F}. The gas associated directly with G10.6$-$0.4~is detected at a systemic source velocity of \\mbox{$v_\\mathrm{LSR}\\approx -1$\\,km\\,s$^{-1}$}, determined from OH maser emission observations, while the foreground gas is detected at \\mbox{$v_\\mathrm{LSR}\\approx 10-55$\\,km\\,s$^{-1}$}.       Section~\\ref{observations} summarises the observations and data reduction, and the results from our \\emph{Herschel} observations are found in Sect.~\\ref{section: results}. The hyperfine structure (hfs) components of the nitrogen hydrides are discussed in Sect.~4. In Sect.~\\ref{Section: Analysis of abundances in different velocity components} we use three different methods to decompose the absorption lines   along the sight-lines towards both sources    in different velocity components, and  estimate  column densities, $N$,  and  relative abundances, $X$,   in each   component. We also compare the column densities of the nitrogen hydrides, both with each other to investigate possible correlations, and with other species tracing regions with  both low and high molecular fractions. Section~\\ref{OPR ammonia} presents our estimates of the ortho-to-para ratio (OPR) of NH$_3$. We end this paper with a summary and outlook in Sect.~\\ref{section summary}. Note that the analysis of the \\emph{background} source emissions and absorptions    is left for a future paper. ",
        "conclusions": "\\label{section summary}   Our spectrally resolved  rotational transitions of NH, o-NH$_2$, ortho- and para-NH$_3$ along the   sight-lines towards the high-mass star-forming regions  W49N  and G10.6$-$0.4 show remarkable similarities of line profiles and abundances. We find  similar abundances of all three species and a co-existence in   diffuse or translucent  interstellar gas with a high molecular fraction. The mean abundance of \\mbox{ortho-NH$_ 3$} abundance is \\mbox{$\\sim$2$\\times$10$^{-9}$} towards both sources. The  mean ratios of all three methods   of the nitrogen hydrides in all velocity components, are  \\mbox{$N$(NH)/$N$(o-NH$_3)$\\,=\\,5.9 and 3.5},  and  \\mbox{$N$(o-NH$_2$)/$N$(o-NH$_3$)\\,=\\,2.4} and 2.0, towards G10.6$-$0.4 and W49N, respectively. This is in sharp contrast to previous observations of the nitrogen hydrides in dark clouds where the ammonia abundances are found to be $\\sim$100 times higher than $X$(NH), and $\\sim$10\\,--\\,300 times higher than $X$(NH$_2$). NH and \\mbox{o-NH$_2$} are  found to be  linearly correlated with  \\mbox{o-NH$_3$} at least for $N$(\\mbox{o-NH$_3$})\\,$\\lesssim$\\,5$\\times$10$^{12}$~cm$^{-2}$ which corresponds to a few $A_\\mathrm{V}$. Upper limits of $N$(NH$^+$) in both sources indicate a $N$(NH$^+$)/$N$(NH) ratio of $\\lesssim$2\\,--\\,14\\,\\%, with a mean of $\\lesssim$6\\,\\%.  Linear correlations are also found for all three nitrogen hydrides with  respect to CH, CN and HNC, although CH displays a  more loose correlation  than the latter two species. The nitrogen hydrides also largely follow the absorption pattern in Doppler velocity space of HCO$^+$ and water, a species also known to trace regions of a high molecular fraction.  We have obtained a surprisingly low  ortho-to-para ratio of ammonia, \\mbox{$\\approx$\\,0.5\\,--\\,0.7$\\pm$0.1}, in the strongest velocity components, which is below the high-temperature limit of unity. No clear explanation  has been found. More observations are needed of both the    rotational  transitions  and the inversion lines with ground-based facilities, to be able to make firm conclusions about the  ammonia OPR in diffuse gas.  We will continue to investigate the absorption lines in the sight-lines towards the other six PRISMAS sources. This will allow an analysis of the nitrogen chemistry at various galactic distances from the Galactic Centre. We will also use new  Open~Time~1 (OT1) \\emph{Herschel}-HIFI data of higher excitation lines to analyse the hot core sources  which will be compared and contrasted  with the diffuse interstellar gas. The ortho-to-para ratio of NH$_3$ will also be further investigated both in  the sources and in the diffuse gas, in addition to the  OPR  of  NH$_2$, for which  new  OT1 data in four of the PRISMAS sources will be analysed and compared to the ammonia OPR."
    },
    "1208/1208.4845_arXiv.txt": {
        "abstract": "We examine model independent constraints on the high redshift and prerecombination expansion history from cosmic microwave background observations, using a combination of principal component analysis and other techniques. This can be translated to model independent limits on early dark energy and the number of relativistic species $\\neff$.  Models such as scaling (Doran-Robbers), dark radiation ($\\Delta\\neff$), and barotropic aether fall into distinct regions of eigenspace and can be easily distinguished from each other. Incoming CMB data will map the expansion history from $z=0$--$10^5$, achieving subpercent precision around recombination, and enable determination of the amount of early dark energy and valuable guidance to its nature. ",
        "introduction": "The expansion history of the universe is a fundamental property of cosmology, reflecting the energy density constituents and their evolution. Yet remarkably little is known in detail about it, other than in a coarse grained average.  For redshifts between 3000 and $10^9$, the universe was mostly radiation dominated, for redshifts between 3000 and $\\sim1$ it was mostly matter dominated, but excursions are possible -- in the effective number of relativistic species $\\neff$ say, or even temporary breakdown of such domination -- and the level of subdominant components is not well constrained. Only around the epoch of primordial nucleosynthesis and of recombination is the expansion rate (Hubble parameter) better constrained, but even there at the $\\sim5\\%$ level averaged over the epoch \\cite{kaplinghat,zahn}.  Given the importance of the expansion history, and the improvement in cosmic microwave background (CMB) data, we investigate what constraints can be placed on it in a model independent way, i.e.\\ other than fitting for a deviation of a particular functional form such as extra $\\neff$ or a specific dark energy model. This would fill in a vast range of cosmic expansion where almost no precise constraints have been placed.  That is, an error band for the Hubble parameter $H(z)$ at $z>1000$ should be a staple of cosmology textbooks, and yet does not exist.  The early expansion history has an important bearing on understanding the nature of dark energy as well, the question of {\\it persistence\\/} of dark energy.  For a cosmological constant $\\Lambda$, the dark energy density contributed at recombination is $\\Omega_\\Lambda\\approx 10^{-9}$, while the current upper limit from data is above $10^{-2}$.  This gives substantial unexplored territory.  Moreover, the current constraints use a specific functional form for the dark energy evolution (usually the Doran-Robbers form \\cite{dorrob}), but other models could lead to significantly different limits \\cite{rdof}.  Thus, model independent limits on early dark energy are needed.  Physics origins for early dark energy can be quite diverse, e.g.\\ from dilaton models (as in some string theories) to k-essence (noncanonical kinetic field theories) to dark radiation (as in some higher dimension theories) \\cite{cope}.  Establishing whether CMB observations could distinguish these classes is another important question.  Improvement of CMB data recently by higher resolution observations extending the temperature power spectrum to multipoles $\\ell\\approx3000$ by the Atacama Cosmology Telescope (ACT \\cite{act}) and South Pole Telescope (SPT \\cite{spt}) gives valuable leverage since higher multipoles are sensitive to modes crossing the cosmological horizon at earlier times. This advance was used in \\cite{linsmith} to rule out in a model independent manner the presence of any epoch of cosmic acceleration between $z\\approx2$ and $10^5$ (supplementing the limits from growth of structure post-recombination in \\cite{unique}).  Upcoming Planck and ground based polarization experiment data will also map out the polarization power spectra, giving additional constraints.  To carry out a model independent analysis of the early expansion history, we use a combination of redshift binning and principal component analysis. In Sec.~\\ref{sec:bins} we lay out the methodology for describing arbitrary $H(z)$.  Analyzing the results in Sec.~\\ref{sec:pca}, we identify the redshifts ranges where the CMB observations are most sensitive to expansion variations.  We project three classes of models representing different physical origins onto the eigenmodes to explore the discriminating power of the data in Sec.~\\ref{sec:edemodels}.  In Sec.~\\ref{sec:concl} we discuss the results and future prospects. ",
        "conclusions": "\\label{sec:concl}  Our knowledge of the expansion history of our universe, even at the level of degree of matter domination or radiation domination at early epochs, is remarkably imprecise.  Cosmic microwave background radiation measurements from ACT, Planck, and SPT (and later ACTpol and SPTpol) will shed light on the times around recombination and reionization. We quantify the model independent state of our knowledge through a combination of redshift bin and principal component analysis, finding that subpercent level constraints will be placed by Planck over $\\log a=[-2.5,-5]$ for a bandwidth of $\\Delta\\log a=0.5$.  CMB data will address one of the key aspects of dark energy -- its persistence, a characteristic of many high energy physics origins -- and we find that several different classes of early dark energy are well separated in PCA space.  The limits can also be interpreted in terms of the number of effective relativistic species, $\\neff$, such as an extra neutrino type, with current data mildly preferring further contributions. A thermal relativistic neutrino species adds 23\\% to the photon energy density, so $\\delta=0.13\\,\\Delta\\neff$, giving tight limits on extra relativistic degrees of freedom from the forthcoming data.  We explore three specific models, representing different classes for early dark energy, possibly corresponding to different physical origins. The commonly used Doran-Robbers form has a dark energy fraction $\\Omega_e$ that is constant through the recombination epoch. We also investigate a dark radiation model with $\\Omega_{de}(a)$ rising to the past and a barotropic aether model with $\\Omega_{de}(a)$ falling to the past, and find that the dominant PC mode is well able to distinguish between these behaviors.  Since the amplitude of that mode is greatest for the Doran-Robbers model, we expect that data constraints on $\\Omega_e$ in the other classes will be weaker than in this model (such as from \\cite{reichardt} using current CMB data), allowing for nonnegligible persistence of dark energy (see \\cite{rdof} for further demonstrations of this). Our general approach, however, does not rely on assuming the form for the new component or expansion deviation.  Figure~\\ref{fig:binFisher} is in a sense the textbook picture of what Planck CMB data will say in a model independent manner about early universe expansion.  For postrecombination epochs this will improve with further ground based polarization measurements (especially of CMB lensing) and inclusion of growth of structure data. Understanding early expansion is in fact crucial for accurate interpretation of large scale structure, and feeds directly into the early time gravitational growth calibration parameter $g_\\star$ \\cite{gstar}; ignorance of this can bias cosmological parameter estimation and tests of gravity.  Expansion history is not the whole story as the effective fluid behind the expansion deviations has perturbations and can have internal degrees of freedom.  We treat the perturbations consistently -- the dark radiation and barotropic aether models for example have sound speeds different from the speed of light.  We do not include viscosity, however, as the data has poor leverage on this \\cite{cold,rdof,smithdaszahn}.  Another difficulty for model independent analysis is having $\\delta<0$, since perturbations are difficult to treat when the effective density deviation passes through zero; models such as nonthermal neutrinos, with energy densities below the standard, could realize such a condition.  We will consider such cases in future work.  Principal component analysis provides a valuable guide to the key epochs of sensitivity and the amount of information contributed from different times.  However, we emphasize and demonstrate that the raw uncertainty $\\sigma_i$ on an eigenmode has very limited meaning; the first 15 modes ordered by $\\sigma_i$ can give a highly inaccurate reconstruction relative to a smaller number of modes ordered by signal to noise.  Redshift bins can be more clearly interpreted.  Employing the best aspects of each can result in physically clear, well characterized expansion history constraints."
    },
    "1208/1208.4132_arXiv.txt": {
        "abstract": "*{This chapter reviews the nature of turbulence in the Galactic interstellar medium (ISM) and its connections to the star formation (SF) process. The ISM is turbulent, magnetized, self-gravitating, and is subject to heating and cooling processes that control its thermodynamic behavior, causing it to behave approximately isobarically, in spite of spanning several orders of magnitude in density and temperature. The turbulence in the warm and hot ionized components of the ISM appears to be trans- or subsonic, and thus to behave nearly incompressibly.  However, the neutral warm and cold components are highly compressible, as a consequence of both thermal instability (TI) in the atomic gas and of moderately-to-strongly supersonic motions in the roughly isothermal cold atomic and molecular components. Within this context, we discuss: i) the production and statistical distribution of turbulent density fluctuations in both isothermal and polytropic media; ii) the nature of the clumps produced by TI, noting that, contrary to classical ideas, they in general accrete mass from their environment in spite of exhibiting sharp discontinuities at their boundaries; iii) the density-magnetic field correlation (and, at low densities, lack thereof) in turbulent density fluctuations, as a consequence of the superposition of the different wave modes in the turbulent flow; iv) the evolution of the mass-to-magnetic flux ratio (MFR) in density fluctuations as they are built up by dynamic compressions; v) the formation of cold, dense clouds aided by TI, in both the hydrodynamic and the magnetohydrodynamic (MHD) cases; vi) the expectation that star-forming molecular clouds are likely to be undergoing global gravitational contraction, rather than being near equilibrium, as generally believed, and vii) the regulation of the star formation rate (SFR) in such gravitationally contracting clouds by stellar feedback which, rather than keeping the clouds from collapsing, evaporates and disperses them while they collapse. }  \\abstract{This chapter reviews the nature of turbulence in the Galactic interstellar medium (ISM) and its connections to the star formation (SF) process. The ISM is turbulent, magnetized, self-gravitating, and is subject to heating and cooling processes that control its thermodynamic behavior, causing it to behave approximately isobarically, in spite of spanning several orders of magnitude in density and temperature. The turbulence in the warm and hot ionized components of the ISM appears to be trans- or subsonic, and thus to behave nearly incompressibly.  However, the neutral warm and cold components are highly compressible, as a consequence of both thermal instability (TI) in the atomic gas and of moderately-to-strongly supersonic motions in the roughly isothermal cold atomic and molecular components. Within this context, we discuss: i) the production and statistical distribution of turbulent density fluctuations in both isothermal and polytropic media; ii) the nature of the clumps produced by TI, noting that, contrary to classical ideas, they in general accrete mass from their environment in spite of exhibiting sharp discontinuities at their boundaries; iii) the density-magnetic field correlation (and, at low densities, lack thereof) in turbulent density fluctuations, as a consequence of the superposition of the different wave modes in the turbulent flow; iv) the evolution of the mass-to-magnetic flux ratio (MFR) in density fluctuations as they are built up by dynamic compressions; v) the formation of cold, dense clouds aided by TI, in both the hydrodynamic (HD) and the magnetohydrodynamic (MHD) cases; vi) the expectation that star-forming molecular clouds are likely to be undergoing global gravitational contraction, rather than being near equilibrium, as generally believed, and vii) the regulation of the star formation rate (SFR) in such gravitationally contracting clouds by stellar feedback which, rather than keeping the clouds from collapsing, evaporates and disperses them while they collapse. } ",
        "introduction": "\\label{sec:intro}  The interstellar medium (ISM) of our galaxy (the Milky Way, or simply, The Galaxy) is mixture of gas, dust, cosmic rays, and magnetic fields that occupy the volume in-between stars \\citep[e.g.,][]{Ferriere01}. The gasesous component, with a total mass $ \\sim 10^{10} \\Msun$, may be in either ionized, neutral atomic or neutral molecular forms, spanning a huge range of densities and temperatures, from the so-called hot ionized medium (HIM), with densities $n \\sim 10^{-2} \\pcc$ and temperatures $T \\sim 10^6$ K, through the warm ionized and neutral (atomic) media (WIM and WNM, respectively, both with $n \\sim 0.3 \\pcc$ and $T \\sim 10^4$ K) and the cold neutral (atomic) medium (CNM, $n \\sim 30 \\pcc$, $T \\sim 100$ K), to the {\\it giant molecular clouds} (GMCs, $n \\gtrsim 100 \\pcc$ and $T \\sim 10$--20 K). These span several tens of parsecs across, and, in turn, contain plenty of substructure, which is commonly classified into {\\it clouds} ($n \\sim 10^3 \\pcc$, size scales $L$ of a few parsecs), {\\it clumps} ($n \\sim 10^4 \\pcc$, $L \\sim 1$ pc), and {\\it cores} ($n \\gtrsim 10^5 \\pcc$, $L \\sim 0.1$ pc). It is worth noting that the temperature of most molecular gas is remarkably uniform, $\\sim 10$--30 K.  Moreover, the ISM is most certainly turbulent, as typical estimates of the Reynolds number ($\\Rey$) within it are very large. For example, in the cold ISM, $R_{\\rm e} \\sim 10^5$--$10^7$ \\citep[][\\S 4.1]{ES04}. This is mostly due to the very large spatial scales involved in interstellar flows.  Because the temperature of the ISM varies so much from one type of region to another, so does the sound speed, and therefore the turbulent velocity fluctuations are often moderately or even strongly supersonic \\citep[e.g., ][and references therein]{HT03, ES04}. In these cases, the flow is significantly compressible, inducing large-amplitude (nonlinear) density fluctuations. The density enhancements thus formed constitute dense clouds and their substructure \\citep[e.g.,][]{Sasao73, Elm93, BP+99a}.  In addition to being turbulent, the ISM is subject to a number of additional physical processes, such as gravitational forces exerted by the stellar and dark matter components as well as by its own self-gravity, magnetic fields, cooling by radiative microscopic processes, and radiative heating due both to nearby stellar sources as well as to diffuse background radiative fields. It is within this complex and dynamical medium that stars are formed by the gravitational collapse of certain gas parcels.  In this chapter, we focus on the interaction between turbulence, the effects of radiative heating and cooling, which effectively enhance the compressibility of the flow, the self-gravity of the gas, and magnetic fields. Their complex interactions have a direct effect on the star formation process. The plan of the chapter is as follows: in \\S \\ref{sec:thermodynamics} we briefly recall the effects that the net heating and cooling have on the effective equation of state of the flow and, in the case of thermally unstable flows, on its tendency to spontaneously segregate in distinct phases. Next, in \\S \\ref{sec:turb} we discuss a few basic notions about turbulence and the turbulent production of density fluctuations in both the hydrodynamic (HD) and magnetohydrodynamic cases, to then discuss, in \\S \\ref{sec:turb_therm}, the evolution and properties of clouds and clumps formed by turbulence in multiphase media. In \\S \\ref{sec:turb_multi}, we discuss the likely nature of turbulence in the diffuse (warm and hot) components of the ISM, as well as in the dense, cold atomic and molecular clouds, suggesting that in the latter, at least during the process of forming stars, the velocity field may be dominated by gravitational contraction. Next, in \\S \\ref{sec:turb_SF} we discuss the regulation of star-formation (SF) in gravitationally contracting molecular clouds (MCs), in particular whether it is accomplished by magnetic support, turbulence, or stellar feedback, and how. Finally, in \\S \\ref{sec:conclusions} we conclude with a summary and some final remarks. ",
        "conclusions": "\\label{sec:conclusions}  In this contribution, we have briefly reviewed the role and interaction between the main physical processes present in the ISM: radiative heating and cooling, magnetic fields, self-gravity, and turbulence, and their implications for the SF process. The presence of radiative heating and cooling implies in general that the gas behaves in a non-isentropic (i.e., non-adiabatic) way, and in particular it may become {\\it thermally unstable} in certain regimes of density and temperature, where low-amplitude (i.e., {\\it linear}) perturbations can cause runaway heating or cooling of the gas that only stops when the gas exits that particular regime. This in turn causes the gas to avoid those unstable density and temperature ranges, and to settle in the stable ones, thus tending to segregate the gas into different phases of different densities and/or temperatures. In classical models of the ISM, only the stable phases were expected to exist in significant amounts.  We then discussed some compressible MHD turbulence basics, and the production, nature and evolution of turbulent density fluctuations in polytropic (i.e., of the form $P \\propto \\rho^\\gamef$) flows, discussing in particular the probability density function (PDF) of the density fluctuations, which takes a lognormal form in isothermal regimes, and develops power-law tails in polytropic ones. We also discussed the correlation (and, at low densities, lack thereof) between the magnetic field and the density as a consequence of the superposition of the different MHD wave modes, and the evolution of the mass-to-magnetic flux ratio (MFR) as density enhancements are assembled by turbulent fluctuations.  We next discussed turbulence in the multi-phase ISM, noting that, since turbulence is an inherently mixing phenomenon, it opposes the segregating effect of thermal instability, causing the production of gas parcels in the classically forbidden unstable regimes, which may add up to nearly half the mass of the ISM, although the density PDF in general still exhibits some multimodality due to the preference of the gas to settle in the stable regimes. The existence of gas in the unstable ranges has been established by various observational studies.  Next, we discussed the nature of the turbulence in the different ranges of density and temperature of the gas, noting that in the diffuse ionized regions, where the flow is transonic (i.e., with Mach numbers $\\Ms \\sim 1$), the gas appears to behave in an essentially incompressible way, exhibiting Kolmogorov scalings over many orders of magnitude in length scale. However, in the neutral atomic component, where the gas is thermally bistable, the flow is expected to exhibit large density and temperature fluctuations, by up to factors $\\sim 100$, thus being highly fragmented. We also pointed out that large-scale compressions in the warm neutral gas, which may be triggered by either random turbulent motions, or by yet larger-scale instabilities, may nonlinearly induce the formation of large regions of dense, cold gas; much larger, in particular, than the most unstable scales of TI, which have sizes $\\sim 0.1$ pc, thus forming large cold atomic clouds that may be the precursors of giant molecular clouds (GMCs). This is because these clouds are expected to become molecular, gravitationally unstable, and magnetically supercritical at approximately the same time, so that when they reach a mostly molecular stage, they are likely to be undergoing generalized gravitational contraction.  The clouds are born internally turbulent and clumpy, and the resulting nonlinear density fluctuations (``clumps'') eventually become locally gravitationally unstable during the contraction of the whole large-scale cloud. Because they are denser, they have shorter free-fall times, and can complete their local collapses before the global one does, thus producing a regime of {\\it hierarchical gravitational fragmentation}, with small-scale, short-timescale collapses occurring within larger-scale, longer-timescale ones. It is thus quite likely that the flow regime in the dense molecular clouds corresponds to a dominant multi-scale gravitational contraction, with smaller-amplitude random (turbulent) motions superposed on it.  The local collapses cause star formation (SF) that begins before the global collapse is concluded, and the ionizing feedback from the massive stars that form during this stage appears to be sufficient to erode and disperse the clouds before the entire mass of the clouds is converted to stars, thus avoiding the objection by \\citet{ZP74} to free-falling GMCs, that they would form stars at a rate much larger than the observed Galactic rate. They do so, but only for short periods of time, before most of their mass gets dispersed.  We conclude that turbulence in the magnetized, multi-phase, self-gravitating ISM is an extremely rich and complex phenomenon, but whose (thermo)dynamics is beginning to be understood, together with its relation to the star formation process.  \\begin{acknowledgement} This work has been funded in part by CONACYT grant 102488. \\end{acknowledgement}"
    },
    "1208/1208.4418_arXiv.txt": {
        "abstract": "{Photonic technologies have received growing consideration for incorporation into next-generation astronomical instrumentation, owing to their miniature footprint and inherent robustness. In this paper we present results from the first on-telescope demonstration of a miniature photonic spectrograph for astronomy, by obtaining spectra spanning the entire H-band from several stellar targets. The prototype was tested on the 3.9~m Anglo-Australian telescope. In particular, we present a spectrum of the variable star $\\pi$ 1 Gru, with observed CO molecular absorption bands, at a resolving power $R = 2500$ at 1600 nm. Furthermore, we successfully demonstrate the simultaneous acquisition of multiple spectra with a single spectrograph chip by using multiple fibre inputs. } {} {} {} {} ",
        "introduction": "An integrated photonic spectrograph (IPS) is a miniaturised, monolithic dispersive device. These wafer-based components are typically only several centimetres in size, which makes them robust against misalignments due to environmental factors. Such properties are highly sought after for use in next-generation astronomical instrumentation, because spectrographs for seeing-limited telescopes, for example, are currently based on large, custom-made optics, which are prone to flexure and thermal drift. The possibility of exploiting IPSs in astronomy was proposed as early as 1995 \\citep{W95}, but it was not until the technology reached maturity that it was considered to be a viable option \\citep{BH2006SPIE}.  Integrated photonic spectrographs are most commonly fabricated via lithographic methods. This essentially involves depositing material onto a substrate through a mask, which is used to outline the required circuit. However, once the mask is created, it is a relatively inexpensive process to mass-produce hundreds or even thousands of miniature spectrographs. This photonic approach to spectrograph design allows for small, mass-fabricated, modular components to be used instead of the large, custom-built elements used in existing spectrographs \\citep{ASBH10,PIMMS10}. Furthermore, photonic spectrographs are commonly designed to operate at the diffraction-limit, circumventing the spectrograph-telescope size relation that often plagues conventional spectrographs on large seeing-limited telescopes \\citep{BH2006SPIE,ASBH10}.  In 2009, we experimentally demonstrated the feasibility of this technology for astronomy by measuring atmospheric OH emission lines in the H-band (1460-1810 nm) using a single photonic chip, an early IPS prototype \\citep{ME09}. More recently, in \\citet{ME12}; hereafter CV12, we characterised an improved prototype IPS that could observe multiple sources simultaneously in a laboratory environment, which serves as the basis for the instrument presented in this work. Here we present the first stellar spectral features obtained using a photonic spectrograph from an on-sky test at the 3.9~m Anglo-Australian Telescope (AAT). ",
        "conclusions": "We have demonstrated that it is possible to interface a diffraction-limited photonic spectrograph to a major seeing-limited research telescope and obtain meaningful astronomical spectra. The overall throughput of this early IPS prototype is  poor when compared to modern spectrographs, which was due to the unoptimised nature of the interface optics, but more significantly to the inherent difficulty in coupling a seeing-limited telescope PSF into a photonic spectrograph. Fortunately, these problems could be overcome with a more optimised telescope interface design. Furthermore, because the majority of telescopes in use (or under construction) employ (or plan to employ) an adaptive optics system, a near diffraction-limited PSF will be available, allowing for efficient coupling directly into an SMF, albeit currently over a very small field of view. Hence, a small-field IFU type IPS system could potentially be used without the need for a photonic lantern with throughputs and resolving powers competitive with modern spectrographs, while maintaining the inherent benefits of an all-photonic platform. Space-based applications would obviously be ideally suited to this photonic approach, but remain untested.  If an IPS system could make use of a near diffraction-limited PSF, and operate with a sufficiently broad free-spectral range (100's of nm), it would make cross-dispersion unnecessary. Therefore, the only non-photonic part of our IPS prototype could be removed entirely, and the detector bonded directly to the photonic chip (or multiple stacked chips). This would also require a new generation of detectors with pixel sizes capable of Nyquist-sampling the chip PSF of $\\sim$~7~$\\mu$m along one axis."
    },
    "1208/1208.6301_arXiv.txt": {
        "abstract": "Dark Energy studies with type Ia supernovae set very tight constraints on the photometric calibration of the imagers used to detect the supernovae and follow up their flux variations. Among the key challenges is the measurement of the shape and normalization of the instrumental throughput. The DICE system was developed by members of the Supernova Legacy Survey (SNLS), building upon the lessons learnt working with the MegaCam imager.  It consists in a very stable light source, placed in the telescope enclosure, and generating compact, conical beams, yielding an almost flat illumination of the imager focal plane.  The calibration light is generated by narrow spectrum LEDs selected to cover the entire wavelength range of the imager. It is monitored in real time using control photodiodes.  A first DICE demonstrator, SnDICE has been installed at CFHT. A second generation instrument (SkyDICE) has been installed in the enclosure of the SkyMapper telescope. We present the main goals of the project.  We discuss the main difficulties encountered when trying to calibrate a wide field imager, such as MegaCam (or SkyMapper) using such a calibrated light source. ",
        "introduction": "Modern precision cosmology, such as the measurement of the Dark Energy equation of state with type Ia supernovae (SNe~Ia) \\cite[see e.g.][and references therein]{sullivan_2011, conley_2011, guy_2010, kessler_2009} sets very tight constraints on the accuracy of the flux calibration of the imagers.  Indeed, the measurement of the luminosity distance of high-redshift (resp. low-redshift) supernovae is primarily performed with the redder (resp. bluer) bands of the imagers. Hence, measuring cosmological parameters with SNe~Ia ultimately boils down to comparing fluxes measured either with red and blue passbands and it is fundamental to control the intercalibration of the imager passbands. Taking full advantage of statistics and quality of SN~Ia measurements requires to control this intercalibration with an accuracy of a fraction of a percent.  The current photometric calibration techniques rely on observations of spectrophotometric stellar calibrators.  Establishing such primary standards is notoriously difficult, as one has to anchor astronomical observations to a physical flux scale. One of the best efforts so far is the work of the CALSPEC team \\cite[][and references therein]{CALSPEC}.  The CALSPEC flux scale relies on NLTE {\\em models} of three (and now five) pure hydrogen white dwarfs.  These stars are the primary standards used to calibrate the flux response of the HST instruments, in particular STIS and NICMOS, which are used in turn to extend the CALSPEC library by adding secondary spectrophotometric standards. The SNLS and SDSS-II surveys have chosen to anchor their flux calibration on this so-called HST white dwarf flux scale \\citep{betoule_2012, regnault_2009, holtzman_2008}. The uncertainty that affects the spectrum of the primary standards is however difficult to assess. As the precision of the calibration efforts improves, it seems increasingly important to check these stellar calibrators using laboratory standards.  The wavelength positioning of the survey passbands has also a sizeable impact on the cosmological parameter measurements, as shown for example in table 9 of \\cite{conley_2011}.  The required wavelength accuracy on the filter cut-offs is as low as a fraction of a nanometer. Passband models are usually derived from pre-installation test-bench measurements of the imager optical components. This is not entirely satisfactory, as filters may evolve over time time as shown for example in \\cite{doi_2010}, and it seems necessary to be able to measure and follow up {\\em in situ} the instrument passbands.  With these requirements in mind, several groups have sought to develop instrumental calibration systems, i.e.  light sources that can illuminate the telescope pupil with well characterized light. By todays standards, ``well characterized'' means that the calibration beam has been mapped using Si photodiodes procured from an institute of standards such as the american Institute of Standards and Technology (NIST).  Many different designs have been proposed over the last few years, and quite a few are now being tested on various wide field imagers.  In what follows, we describe the DICE system. DICE stands for Direct Illumination Calibration Experiment. It consists in a very stable point-like source, generating conical beams that deliver a quasi-uniform illumination on the focal plane. Two such systems have been built so far. A first prototype was installed at the Canada France Hawaii Telescope (CFHT) in order to calibrate the 1 deg$^2$ MegaCam imager. A second generation demonstrator was recently installed in the enclosure of the 5.7 deg$^2$ SkyMapper imager. We discuss below the main design (\\S \\ref{sec:design_considerations}) and implementation (\\S \\ref{sec:dice_system}) aspects of the project. We then describe the test bench procedures that permit to characterize the light source (\\S \\ref{sec:test_bench}). Finally, we discuss a few key problems that arise in the data analysis (\\S \\ref{sec:analysis}). ",
        "conclusions": "Instrumental calibration is today a very active subject. Many different light source projects have been proposed and it is quite difficult to predict which design is going to prevail in future surveys. DICE is one of those attempts. It consists in a compact, versatile, inexpensive LED based calibrated light source, that can be placed in the dome of virtually any telescope.  As a point like source placed at a finite distance, it generates a conical beam that yields a quasi-uniform focal plane illumination. This is complemented by a pencil beam that allows (1) to control the relative positions and orientations of the source and the telescope and (2) to study in detail the ghost contamination.  The stability of the light source, permits to carry out a daily monitoring of the telescope and imager response.  The simplicity of the calibration beams, combined with the complementarity pencil beam / conical beam allows to derive an accurate estimate of the telescope throughput, taking into account the contamination by stray light.  This project does not address the estimation of the atmospheric transmission, which is an important (and in some bands, highly variable) component of the telescope effective throughput. Many contributions to these proceedings are dealing with this subject."
    },
    "1208/1208.3712_arXiv.txt": {
        "abstract": "We discuss the discovery and characterization of the circumbinary planet Kepler-38b.  The stellar binary is single-lined, with a period of 18.8 days, and consists of a moderately  evolved main-sequence star ($M_A=0.949\\pm 0.059\\,M_{\\odot}$ and $R_A=1.757\\pm 0.034\\,R_{\\odot}$) paired with a low-mass star ($M_B=0.249\\pm 0.010\\,M_{\\odot}$ and $R_B=0.2724\\pm 0.0053\\,R_{\\odot}$) in a mildly eccentric ($e=0.103$) orbit. A total of eight transits due to a circumbinary planet crossing the primary star were identified in the {\\em Kepler} light curve (using {\\em Kepler} Quarters 1 through 11), from which a planetary period of $105.595\\pm 0.053$ days can be established. A photometric dynamical model fit to the radial velocity curve and {\\em Kepler} light curve yields a planetary radius of $4.35\\pm 0.11\\,R_{\\oplus}$, or $1.12\\pm 0.03\\,R_{\\rm Nep}$.  Since the planet is not sufficiently massive to observably alter the orbit of the binary from Keplerian motion, we can only place an upper limit on the mass of the planet of $122\\,M_{\\oplus}$ ($7.11\\,M_{\\rm Nep}$ or $0.384\\,M_{\\rm Jup}$) at 95\\% confidence. This upper limit should decrease as more {\\em Kepler} data become available. ",
        "introduction": "While the {\\it Kepler} Mission (Borucki et al.\\ 2010) is sometimes considered synonymous with ``the search for Earth-like planets'', its goals are considerably broader, and include estimating the frequency and orbital distribution of planets in multiple-stellar systems. To achieve its goals, {\\it Kepler} relies on its exquisite photometric precision, its ability to simultaneously observe roughly 160,000 stars, and its long-duration and near-continuous time series measurements (Koch et al.\\ 2010). This triad of unique capabilities makes {\\it Kepler} ideally suited for exoplanet discovery and characterization, including planets in binary star systems [see Haghighipour (2010) for an in-depth discussions of planets in binary star systems].   If the binary star's orbital plane is favorably oriented, the stars will eclipse and thus reveal their binary nature. {\\it Kepler} has discovered over 2000 eclipsing binaries (Pr\\v sa et al.\\ 2011; Slawson et al.\\ 2011), with periods ranging from 0.075 to over 275 days, and these systems are being searched for the presence of planets. The eclipses tell us that we are viewing the binary system close to its orbital plane, and thus perhaps at a favorable orientation for finding transiting planets if the planets lie in the same orbital plane. However, detecting such planets is much more difficult than finding planets orbiting a single star. A dilution factor is present, but the main challenges arise from the fact that the transits are neither periodic nor equal in duration (e.g.\\ see Doyle et al.\\ 2011; Welsh et al.\\ 2012). In addition, the deep stellar eclipses can easily mask a small transit signal. Partially compensating for these disadvantages, the timing of the eclipses of the binary component stars provides a very sensitive indicator of the presence of a third body in the system (e.g.\\ Orosz et al.\\ 2012). The eclipse timing variations (ETVs) as seen in an O-C diagram (Observed-minus-Computed) can reveal deviations from periodicity that are attributed to a gravitational perturbation caused by a planet. Note that for short orbital-period binaries, the ETVs are generally dominated by dynamical effects, not light-travel time delays.  The first transiting circumbinary planet discovered was Kepler-16b (Doyle et al.\\ 2011). The transits left no room for ambiguity as to the planetary nature of the third object. The planet is in a P-type orbit (Dvorak 1984, 1986), meaning the planet is circumbinary (an outer orbit around both stars). Soon after, two more transiting circumbinary planets were discovered, Kepler-34b and Kepler-35b (Welsh et al.\\ 2012), establishing that such planets are not rare. While there is considerable diversity among the three systems (in mass ratios, eccentricities, orbital periods), two features are in common: (i) all three planets have a radius similar to Saturn's, which is interesting in that Jupiter-radius planets should be easier to detect; and (ii) the orbital periods of the planets are only slightly longer than the minimum needed to guarantee dynamical stability according to the criteria given in Holman \\& Wiegert (1999). Whether this is a consequence of planet formation and migration, or simply a selection effect, is unknown.  In this paper we announce the discovery of a fourth transiting circumbinary planet, Kepler-38b. As with the other cases, the detection was made by visual inspection of a subset of the eclipsing binary star light curves, namely, those with orbital period greater than $\\sim$1 day. The observations are presented in \\S2, and the photometric-dynamical model fit in \\S3.  We conclude with a discussion in \\S4. ",
        "conclusions": "\\subsection{Planetary Parameters}   The radius of Kepler-38b is $4.35\\,R_{\\oplus}$ ($=1.12\\,R_{\\rm Nep}$ or $0.39\\,R_{\\rm Jup}$, using the equatorial radii), with an uncertainty of $\\pm 0.11\\,R_{\\oplus}$ (or 2.5\\%). For comparison, its radius is about half of the radius of Kepler-16b ($R=0.7538 \\pm 0.0025\\,R_{\\rm Jup}$, Doyle et al.\\ 2011), Kepler-34b ($R=0.764\\pm0.014\\,R_{\\rm Jup}$, Welsh et al.\\ 2012), and Kepler 35b  ($R=0.728\\pm 0.014\\,R_{\\rm Jup}$, Welsh et al.\\ 2012).  Thus all four of the transiting circumbinary planets discovered so far have radii substantially smaller than Jupiter's.  Since a Jupiter-sized planet would have a deeper transit and would therefore be easier to find (all other conditions being equal), the tendency for the circumbinary planets to be sub-Jupiter size is noteworthy. \\citet{Pierens_2008} argued that Jupiter-{\\em mass} circumbinary planets in orbits relatively close to the binary should be rare owing to the various instabilities that occur during the migration phase and also to subsequent resonant interactions with the binary (Jupiter-mass circumbinary planets that orbit further out from the binary could be stable, but these would be less likely to transit owing to a larger separation). The predictions of \\citet{Pierens_2008} seem to be consistent with what is known from the first four {\\em Kepler} transiting circumbinary planets.  Since Kepler-38b has not yet noticeably perturbed the stellar orbits, we have only an upper limit on its mass of $M_b<122\\,M_{\\oplus}$ ($<7.11\\,M_{\\rm Nep}$ or $<0.384\\,M_{\\rm Jup}$) at 95\\% confidence.  While this is clearly a substellar mass, this upper limit is not particularly constraining in terms of the density, as we  find $\\rho_{\\rm b}<8.18$ g cm$^{-3}$. A reasonable mass is $M_b\\approx 21\\,M_{\\oplus}$, assuming the planet follows the empirical mass-radius relation of $M_b=(R_b/R_{\\oplus})^{2.06}\\,M_{\\oplus}$ (Lissauer et al.\\ 2011).  The gravitational interaction between the planet and the two stars causes small perturbations that will grow over time and will eventually lead to a measurable change in the phase difference between the primary and secondary eclipses.  As discussed above, this change in the phase difference manifests itself as a difference between the period measured from the primary eclipses and the period measured from the secondary eclipses. In Kepler-16, Kepler-34, and Kepler-35, the time-scale for a measurable period difference to occur is relatively short, as divergent periods were measured using 6 Quarters of data for Kepler-16, and 9 Quarters of data for Kepler-34 and Kepler-35.  For Kepler-38, the $1\\sigma$ limit on the period difference is $<4.4$ seconds using 11 Quarters. Figure \\ref{dp_vs_M} shows the set of acceptable planetary masses and the changes in orbital period the planet would induce, based on the Monte Carlo Markov Chain from the photometric dynamical model. The larger the planetary mass, the larger the difference in periods it causes between the primary and secondary star. The lack of any measurable period difference places an upper limit on the mass of the planet: To the left of the vertical dashed line at 122 $M_{\\oplus}$ is where 95\\% of the acceptable solutions reside. The vertical dot-dash line marks a planetary mass of 21 $M_{\\oplus}$. The horizontal dashed line at 4.4 seconds marks the observed $1\\sigma$ uncertainty in the measured value of $P_{2} - P_{1}$ accumulated over the span of the current {\\it Kepler} observations (52 binary star eclipses); valid planetary masses lie below this line, roughly. The dotted horizontal line at 0.9 seconds marks a $1\\sigma$ period difference uncertainty that can be placed when $\\approx 150$ eclipses are eventually observed by {\\it Kepler}, corresponding approximately to the end of the Extended Mission in the year 2017. If no period difference is measured at that time, the mass of the planet would be less than $\\approx  20\\,M_{\\oplus}$ with $1\\sigma$ uncertainty. (The intersection of this period difference {\\em uncertainty} and the 21 $M_{\\oplus}$  line within the set of Monte Carlo points is a coincidence.) If the planet has a normal density and a mass of $21\\,M_{\\oplus}$, then the period difference would be $\\approx 0.9$ seconds, which would require $\\approx 312$ binary orbits or about 16 years to obtain a $3\\sigma$ detection.  \\subsection{Stellar Parameters and Age}  As noted earlier, the observed timing of the planet transits and the amplitude of the radial velocity curve sets the scale of the binary, and we are able to measure masses and radii for each star.  Considering the fact that Kepler-38 is a single-lined spectroscopic binary (note that the ratio of the secondary-to-primary flux in the {\\em Kepler} bandpass is $9\\times 10^{-4}$, see Table \\ref{photodynamicalparm}), the uncertainties in the masses and radii are fairly small (6.2\\% and 3.6\\% for the primary and secondary  masses, respectively, and 1.8\\% for the radii).   On the other hand, these uncertainties are still somewhat larger than what one would like when doing precise comparisons with stellar evolutionary models (for comparison, the stellar masses and radii are known to much better than 1\\% for the first three {\\em Kepler} circumbinary planets, see Doyle et al.\\ 2011 and Welsh et al.\\ 2012). Clearly, since the radius of the primary star ($1.757\\pm 0.034\\,R_{\\odot}$) is much larger than the expected zero-age main sequence radius for its mass ($0.949\\pm 0.059\\,M_{\\odot}$), the primary must be significantly evolved, but is still a core hydrogen-burning star. The situation is shown in Figure \\ref{iso}, which gives the position of Kepler-38 in a $T_{\\rm eff}-\\log g$ diagram. The heavy solid line is a Yonsei-Yale evolutionary track (Yi et al.\\ 2001) for $[{\\rm Fe/H}]=-0.11$, which is our adopted spectroscopic metallicity determination (Table 4). As indicated earlier, we assume here that the iron abundance is well approximated by the [m/H] metallicity index measured with SPC. The dark shaded area is the uncertainty in the location of the track that results from the uncertainty in the mass, and the lighter shaded area also includes the uncertainty in the metallicity.  The observed $T_{\\rm eff}$ and $\\log g$ is just outside the $1\\sigma$ region of the evolutionary track. Figure \\ref{massrad} shows the positions of the stars in Kepler-38 on the mass-radius and mass-temperature diagrams [the temperature of the secondary was inferred from the measured temperature of the primary (Table \\ref{erik}) and the temperature ratio from the ELC models (Table \\ref{ELCparm})]. They are compared against model isochrones from the Dartmouth series (Dotter et al. 2008), in which the physical ingredients such as the equation of state and the boundary conditions are designed to better approximate low-mass stars (see, e.g., Feiden et al.\\ 2011). If we use the mass, radius, and metallicity of the primary, the inferred age is 13 Gyr.  If, on the other hand, we use the measured temperature, the age of the system would be between 7 and 8 Gyr. Accounting for the small discrepancy between the measured temperature and the other parameters, we adopt an age of $10\\pm 3$ Gyr.   The secondary star has a relatively low mass ($0.249\\pm 0.010\\,M_{\\odot}$), and is one of just a handful of low mass stars with a well-measured mass and radius. Its mass is slightly larger than those of CM~Dra~A ($0.2310\\, M_{\\odot}$; Morales et al.\\ 2009) and KOI-126~B ($0.20133\\,M_{\\odot}$; Carter, et al.\\ 2011), and Kepler-16b ($0.20255\\pm 0.00066\\,M_{\\odot}$; Doyle et al.\\ 2011, see also Bender et al.\\ 2012 and Winn et al.\\ 2011). Stars whose mass is $<0.8\\,M_{\\odot}$ typically have radii that are $\\sim 10-15\\%$ larger than what is predicted by stellar evolutionary models (Torres \\& Ribas 2002; Ribas 2006; Ribas et al.\\ 2006; L\\'opez-Morales 2007; Torres et al.\\ 2010; Feiden et al.\\ 2011).   Relatively high levels of stellar activity induced by tidal interactions in short-period systems is one possible cause of this discrepancy (L\\'opez-Morales 2007).  There is a hint that the secondary star is inflated, although we note there is a small discrepancy with the models for the primary.    \\subsection{Stellar Variability}  With {\\it Kepler} data, it is frequently possible to determine the rotation period of the star from the modulations in the light curve due to star-spots. In the case of Kepler-38 however, the modulations are small (rms $<600$ ppm for the long cadence time series, omitting the eclipses), and considerably smaller than the instrumental systematic trends in the light curve. Separating intrinsic stellar variability from instrumental artifacts is difficult, so we used the PDC light curve that has many of the instrumental trends removed. The PDC data for Kepler-38 are generally flat outside of eclipses, with the exception of Quarter 1, which we therefore omitted. The data were normalized and the primary and secondary eclipses removed from the time series. A power spectrum/periodogram was computed, and the dominant frequency present corresponds to the orbital period (18.79 days). We also patched the gaps in the light curve with a random walk and computed the auto-correlation function (ACF). The ACF revealed a broad peak at $\\sim 18$ days, consistent with the orbital period.  The fact that the orbital period is manifest in the out-of-eclipse light curve initially suggested that star-spots were present and the star's spin was tidally locked with the orbital period. But the non-zero eccentricity of the orbit means that exact synchronicity is impossible (see below). Phase-folding the data on the orbital period, and then binning to reduce the noise, revealed the origin of the orbital modulation: Doppler boosting combined with reflection and ellipsoidal modulations. The amplitude is $\\sim 300$ ppm with a maximum near phase $\\phi=0.25$, consistent with what is expected from the ELC models. We verified that the Doppler boost signal is also present in the SAP light curves. For more discussion of Doppler boosting in {\\em Kepler} light curves, see Faigler \\& Mazeh (2011) and Shporer et al.\\ (2011).  For an eccentric orbit, there is no one spin period that can be synchronous over the entire orbit. However, there is a spin period such that, integrated over an orbit, there is no {\\em net} torque on the star's spin caused by the companion star. At this ``pseudosynchronous'' period, the spin is in equilibrium and will not evolve (Hut 1981, 1982). Given the old age of the binary, it should have reached this pseudosynchronous equilibrium state. The ratio of orbital period to pseudosynchronous spin period is a function of the eccentricity only (Hut 1981), and for Kepler-38 this period is $P_{\\rm pseudo}=17.7$ days. This is close to, but slightly shorter than, the $\\sim 18.79$ day orbital period. Using this pseudosynchronous spin period and the measured stellar radius, a projected rotational velocity of $V_{\\rm rot} \\sin{i}\\approx 4.7$ km s$^{-1}$ is expected, if the star's spin axis is aligned with the binary orbital axis. Our spectral modeling yields $V_{\\rm rot} \\sin{i} = 2.4 \\pm 0.5$ km s$^{-1}$, which is close to the rotational velocity expected due to pseudosynchronous rotation. Given that the rotational velocity is near the spectral resolution limit, we can't rule out systematic errors of a few km s$^{-1}$ caused by changes in the instrumental point-spread function, macroturbulence, etc., which could bring the measured value of $V_{\\rm rot}\\sin i$ up to the pseudosynchronous value.   \\subsection{Orbital Stability and the Habitable Zone}   The 105 day planetary orbital period is the shortest among the first four {\\em Kepler} circumbinary planets. The planet orbits the binary quite closely: the ratio of planetary to stellar orbital periods is 5.6 (ratio of semi-major axes is 3.2). Such a tight orbit is subject to dynamical perturbations, and following the analytic approximation given by Holman \\& Wiegert (1999), the critical orbital period in this binary below which the planet's orbit could experience an instability is 81 days. This compares favorably with the results of direct $N$-body integrations, which yield a critical orbit period of 75 days. Thus while stable, the planet is only 42\\% above the critical period (or 26\\% beyond the critical semi-major axis). Kepler-38 thus joins Kepler-16 (14\\%), Kepler-34 (21\\%), and Kepler-35 (24\\%) as systems where the planet's orbital period is only modestly larger than the threshold for stability. The fact that the first four circumbinary planets detected by {\\em Kepler} are close to the inner stability limit is an interesting orbital feature that may be explained by processes such as planetary migration and planet-planet scattering during and/or post formation of these objects. For example, it is also possible that strong planet migration will  bring planets close in: migration may cease near the instability separation leading to a pile-up just outside the critical radius; or planets that continue to migrate in are dynamically ejected, leaving only those outside the critical radius \\citep{Pierens_2008}. There is also an observational bias since objects that orbit closer to their host star(s) will be more likely to transit, making them more likely to be discovered.  Regardless of the cause, the close-to-critical orbits have an interesting consequence. Many {\\em Kepler} eclipsing binaries have orbital periods in the range 15-50 days and have G and K type stars \\citep{Prsa_2011,Slawson_2011}, and for such binaries the critical separation (roughly 2-4 times the binary separation, depending on the eccentricity of the binary) is close to the habitable zone{\\footnote{For a binary star system, the habitable zone is no longer a spherical shell but a more complex shape that rotates with the binary.}}. Thus, observed circumbinary planets may preferentially lie close to their habitable zones. Kepler-16b is just slightly exterior to its habitable zone, while Kepler-34b is slightly interior (too hot). Kepler-38b is well interior to its habitable zone, with a mean equilibrium temperature of $T_{\\rm eq}=475$ K, assuming a Bond albedo of 0.34 (similar to that of Jupiter and Saturn). Although $T_{\\rm eq}$ is somewhat insensitive to the albedo, this temperature estimation neglects the atmosphere of the planet, and therefore should be considered a lower limit. It is interesting to consider the situation at a much earlier time in the past when the primary was near the zero-age main sequence. The primary's luminosity would have a factor $\\approx 3$ smaller, and the equilibrium temperature of the planet would have been $T_{\\rm eq}\\approx 361$~K, assuming a similar orbit and Bond albedo."
    },
    "1208/1208.0711_arXiv.txt": {
        "abstract": "{ Recently we have studied the Lorentzian version of the IIB matrix model as a nonperturbative formulation of superstring theory. By Monte Carlo simulation, we have shown that the notion of time ---as well as space---emerges dynamically from this model, and that we can \\emph{uniquely} extract the real-time dynamics, which turned out to be rather surprising: after some ``critical time'', the SO(9) rotational symmetry of the nine-dimensional space is spontaneously broken down to SO(3) and the three-dimensional space starts to expand rapidly. In this paper, we study the same model based on the classical equations of motion, which are expected to be valid at later times. After providing a general prescription to solve the equations, we examine a class of solutions, which correspond to manifestly commutative space. In particular, we find a solution with an expanding behavior that naturally solves the cosmological constant problem. } ",
        "introduction": "\\label{sec:introduction}   There are many fundamental questions in cosmology, which can, in principle, be answered by superstring theory. Describing the birth of the universe is one of the most fundamental ones. It is recognized, however, that the cosmic singularity is not resolved generally in perturbative string theory \\cite{Liu:2002ft,Liu:2002kb,Lawrence:2002aj,% Horowitz:2002mw,Berkooz:2002je}. Therefore, in order to study the very early universe, we definitely need a nonperturbative formulation. Among various proposals \\cite{BFSS,IKKT,DVV} based on matrix models,\\footnote{For earlier attempts to apply matrix models to cosmology, see ref.\\ \\cite{Freedman:2004xg,Craps:2005wd,Li:2005sz,% Das:2005vd,Chen:2005mga,She:2005mt,Martinec:2006ak,Ishino:2006nx,% Matsuo:2008yd,Klammer:2009ku,Lee:2010zf}.} the IIB matrix model \\cite{IKKT} looks most natural for describing the birth of the universe, since not only space but also time is expected to appear dynamically from the matrix degrees of freedom. Also the model is unique in that it is a manifestly covariant formulation, while the other proposals are based on the light-cone formulation, which breaks covariance.  One of the important issues in the IIB matrix model is to identify the configurations of matrices that dominate the path integral and to determine the corresponding space-time structure. This was studied by various approaches \\cite{AIKKT,Hotta:1998en,Ambjorn:2000bf,% Ambjorn:2000dx,Anagnostopoulos:2001yb,Anagnostopoulos:2011cn,% Nishimura:2000ds,Nishimura:2000wf,Nishimura:2001sx,Kawai:2002jk,Aoyama:2006rk,% Imai:2003vr,Imai:2003jb,Imai:2003ja,Nishimura:2011xy} in the Euclidean version of the model, which was shown to have finite partition function without any cutoffs \\cite{Krauth:1998xh,Austing:2001pk}. However, the Euclidean model does not seem suitable for cosmology since it does not provide the real-time dynamics. Furthermore, a recent study based on the gaussian expansion method suggests that the space-time obtained dynamically in the Euclidean model is three-dimensional rather than four-dimensional \\cite{Nishimura:2011xy}. While this conclusion itself may have a profound implication, we do not know yet how it should be physically interpreted.   All these considerations led us to study the Lorentzian version of the IIB matrix model nonperturbatively \\cite{KNT}. By Monte Carlo simulation, we have shown that the Lorentzian model can be made well-defined nonperturbatively by first introducing infrared cutoffs and then removing them appropriately in the large-$N$ limit. We have also found that the eigenvalue distribution of the matrix in the temporal direction extends in that limit, which implies that time emerges dynamically. (Supersymmetry of the model plays a crucial role.) Indeed we were able to extract a unique real-time dynamics, which turns out to have a surprising property. After some critical time, the SO(9) rotational symmetry of the space is broken spontaneously down to SO(3), and the three-dimensional space starts to expand rapidly. This result can be interpreted as the birth of the universe. Note that the concept of time-evolution also emerges dynamically in this model without ever having to specify the initial condition.  Although the length of time-evolution that has been extracted from Monte Carlo simulation is restricted due to finite $N$, we consider that the whole history of the universe can be obtained from the same model in the large-$N$ limit. If this is true, we should be able to answer various important questions in both cosmology and particle physics. For instance, we may obtain the microscopic description of inflation. If we can reach the time at which stringy excitations and quantum gravitational effects become negligible, we may see how particles in the Standard Model (or possibly its extension) starts to appear. If we can study the behaviors of the model at much later times, we may be able to understand why the expansion of our universe is accelerating in the present epoch. Finally we may even predict how our universe will be like in the future.   While the late-time behaviors are difficult to study by direct Monte Carlo methods, the classical equations of motion are expected to become more and more valid at later times since the value of the action increases with the cosmic expansion. We will see that there are actually many classical solutions, which is reminiscent of the fact that string theory possesses infinitely many vacua that are perturbatively stable. However, unlike in perturbative string theory, we have the possibility to pick up the unique solution that describes our universe by requiring smooth connection to the behavior at earlier times accessible by Monte Carlo simulation. From this perspective, we consider it important to classify the classical solutions and to examine their cosmological implications. The aim of this paper is to make a first step in that direction. In particular, we find a classical solution with an expanding behavior that can naturally solve the cosmological constant problem.   Another issue we would like to address in this paper concerns how the commutative space-time appears from this model. This is important since the SO(9) symmetry breaking observed in the Monte Carlo simulation is understood intuitively by a mechanism, which relies crucially on the fact that the matrices that represent the space-time are noncommutative \\cite{KNT}. We show that classical solutions which correspond to manifestly commutative space can be easily constructed, and discuss the vanishing of noncommutativity between space and time in some simple examples. This implies that the emergence of commutative space-time is indeed possible in this model at later times.     The rest of this paper is organized as follows. In section \\ref{sec:IIBmm} we briefly review the IIB matrix model and the results obtained by our previous Monte Carlo studies \\cite{KNT}. In section \\ref{sec:classical-sol} we provide a general prescription to find classical solutions. In particular, we show that this can be done systematically by using Lie algebras. Then we restrict ourselves to manifestly space-space commutative solutions and present a complete classification of such solutions within certain simplifying ansatz. The particular solution discussed in our previous publication \\cite{KNT2} also appears in this classification. In section \\ref{sec:su11-su2} we obtain explicit time-evolution of the scale factor and the Hubble parameter for some simple solutions, and discuss their cosmological implications. Section \\ref{sec:concl} is devoted to a summary and discussions. In appendix \\ref{sec:unitary-rep} we review the irreducible unitary representations of the SU$(1,1)$ algebra, which will be needed in constructing the solutions discussed in section \\ref{sec:su11-su2}. In appendix \\ref{sec:others} we list the other simple solutions which are not discussed in section \\ref{sec:su11-su2}. In appendix \\ref{sec:not-ss-com} we give some examples of solutions that are not manifestly space-space commutative. ",
        "conclusions": "\\label{sec:concl}  In this paper we studied the late time behaviors of the universe in the Lorentzian version of the IIB matrix model. We investigated the classical equations of motion, which are expected to be valid at later times. This is a complementary approach to Monte Carlo simulation,\\footnote{See ref.~\\cite{Nishimura:2012xs} for a recent review on Monte Carlo studies of matrix models and supersymmetric gauge theories in the context of string theory.} which was used previously to study the birth of the universe in the same model. First we provided a general prescription to solve the equations of motion. The problem reduces to that of finding a unitary representation of a Lie algebra.  In this way, we obtained a class of solutions that are manifestly space-space commutative. The simplest ones in this class are the $d=1$ solutions with $A_2=\\cdots=A_9=0$, from which we can easily construct the ones representing higher dimensional space-time as well. We made a complete classification of such solutions. Some solutions represent expanding $(3+1)$-dimensional universe without space-time noncommutativity in the continuum limit. In particular, we find that there exists a solution, in which the parameter $w$ changes smoothly from $-1$ to $-1/3$. This explains why we seem to have a tiny cosmological constant in the present epoch, and hence can naturally solve the cosmological constant problem. While we do not insist that this particular solution really describes our universe, we consider that the cosmological constant problem can be naturally solved in the Lorentzian matrix model in a similar manner.    Corresponding to what we have done in Monte Carlo simulation, we have introduced infrared cutoffs in both the temporal and spatial directions. These are represented by the Lagrange multipliers $\\lambda$ and $\\tilde{\\lambda}$ introduced in the action (\\ref{tildeS}). In general, this breaks the SO(9,1) symmetry of the model explicitly. Let us note, however, that it is possible to have $\\lambda = \\tilde{\\lambda}$ in the cases (a) and (b) in section \\ref{sec:d1}. Such solutions break the SO(9,1) symmetry \\emph{spontaneously}. We expect that the explicit breaking of the Lorentz symmetry by the infrared cutoffs disappears in the large-$N$ limit. If that is really the case, we should select a solution with $\\lambda = \\tilde{\\lambda}$. It is intriguing to note that the cases (a) and (b) are indeed the ones that are physically interesting.   The (3+1)-dimensional space-time represented by the solutions discussed mainly in this paper has the topology $R\\times S^3$. This is a restriction which we have as long as we construct such solutions based on the $d=1$ solution. In other constructions, we can also obtain solutions representing a space with the topology of a three-dimensional ball as we discussed in appendix \\ref{sec:others}. While the Monte Carlo results seem to be more consistent with the latter topology of the space, it remains to be seen what kind of topology is actually realized at later times.   Below we list some directions for future investigations.  First we consider it important to examine the stability of the solutions we found in this paper. It would be also interesting to calculate the one-loop effective action around the solutions. That would tell us the validity of the solutions, and we should be able to know how late the time should be for the solutions to be valid.   Secondly it is important to understand better how one should extract the information of the space-time metric from a matrix configuration. Ref.~\\cite{Hanada:2005vr} shows that this is indeed possible, in principle, if one interprets the matrix as a covariant derivative on the space-time manifold, where the general coordinate invariance is realized manifestly as a subgroup of the SU($N$) symmetry. However, this interpretation is different from the one adopted in this paper, which is compatible with the supersymmetry as we reviewed in section \\ref{sec:IIBmm}. The precise relationship between the two interpretations is yet to be clarified, although it is tempting to consider that they are related to each other by T-duality of type IIB superstring theory. In this work we have naively identified $R(t)$ with the scale factor in the Friedman-Robertson-Walker metric when we discuss cosmological implications in section \\ref{sec:su11-su2}. It remains to be seen whether this identification can somehow be justified.      Thirdly we consider it important to study a wider class of solutions using the general prescription provided in this paper. In particular, it would be interesting to examine the solutions, which are not manifestly space-space commutative, based on the Lie algebra (\\ref{SO(2,2) and SO(4) solutions}) or the one given in appendix \\ref{sec:not-ss-com}. Also it would be interesting to investigate solutions with nontrivial structure in the extra dimensions. Such structure is expected to play a crucial role \\cite{Chatzistavrakidis:2011gs,Aoki:2010gv} in determining the matter content at late times and in finding how the standard model appears from the matrix model. Eventually, we have to single out the solution, which is smoothly connected to the unique result at earlier times accessible by Monte Carlo simulation.  Developments in the above directions would enable us to solve various fundamental problems in particle physics and cosmology. For instance, we should be able to understand the mechanism of inflation and to clarify what the dark matter and the dark energy are. We hope that the present work will trigger such developments."
    },
    "1208/1208.0182_arXiv.txt": {
        "abstract": "{The Large Magellanic Cloud (LMC) is an ideal target for the study of an unbiased and complete sample of supernova remnants (SNRs). We started an X-ray survey of the LMC with \\xmm, which, in combination with observations at other wavelengths, will allow us to discover and study remnants that are either even fainter or more evolved (or both) than previously known.} {We present new X-ray and radio data of the LMC SNR candidate \\dem, obtained by \\xmm\\ and ATCA, along with archival optical and infrared observations.} {We use data at various wavelengths to study this object and its complex neighbourhood, in particular in the context of the star formation activity, past and present, around the source. We analyse the X-ray spectrum to derive some remnant's properties, such as age and explosion energy.} {Supernova remnant features are detected at all observed wavelengths\\,: soft and extended X-ray emission is observed, arising from a thermal plasma with a temperature $kT$ between 0.2 keV and 0.3 keV. Optical line emission is characterised by an enhanced [\\ion{S}{ii}]-to-H$\\alpha$ ratio and a shell-like morphology, correlating with the X-ray emission. The source is not or only tentatively detected at near-infrared wavelengths (shorter than 10 $\\mu$m), but there is a detection of arc-like emission at mid and far-infrared wavelengths (24 and 70 $\\mu$m) that can be unambiguously associated with the remnant. We suggest that thermal emission from dust heated by stellar radiation and shock waves is the main contributor to the infrared emission. Finally, an extended and faint non-thermal radio emission correlates with the remnant at other wavelengths and we find a radio spectral index between $-$0.7 and $-$0.9, within the range for SNRs. The size of the remnant is $\\sim 79 \\times 64$ pc and we estimate a dynamical age of about 35\\,000 years.} {We definitely confirm \\dem\\ as a new SNR. This object ranks amongst the largest remnants known in the LMC. The numerous massive stars and the recent outburst in star formation around the source strongly suggest that a core-collapse supernova is the progenitor of this remnant.} ",
        "introduction": "\\label{introduction}  Supernova remnants (SNRs) are an important class of objects, as they contribute to the energy balance and chemical enrichment and mixing of the interstellar medium (ISM). However, in our own Galaxy, distance uncertainties and high absorption inhibit the construction of a complete and unbiased sample of SNRs. On the other hand, the Large Magellanic Cloud (LMC) offers a target with a low foreground absorption at a relatively small distance of $\\sim$ 50 kpc \\citep{2008MNRAS.390.1762D}.    Furthermore, the broad multi-frequency coverage of the LMC, from radio to X-rays, allows for easier detection and classification of SNRs, which is most usually done using three signatures\\,: thermal X-ray emission in the (0.2--2~keV) band, optical line emission with enhanced [\\ion{S}{ii}] to H$\\alpha$ ratio \\citep[$\\gtrsim$ 0.4,][]{1973ApJ...180..725M}, and non-thermal (synchrotron) radio-continuum emission, with a typical spectral index of $\\alpha \\sim - 0.5$ (using $S \\propto \\nu\\,^{\\alpha}$, where $S$ is the flux density and $\\nu$ the frequency), although $\\alpha$ can have a wide range of values \\citep{1998A&AS..130..421F}. Nevertheless, the interstellar environment in which the supernova (SN) exploded strongly affects the subsequent evolution of the remnant, so that some SNRs do not exhibit all three conventional signatures simultaneously \\citep{1997AJ....113.1815C}.  \\begin{figure}[!b] \\centering \\includegraphics[width=\\hsize]{N51_paper_label_low.ps} \\caption{The giant \\ion{H}{ii} complex \\mbox{\\object{LHA 120-N 51}} in the light of  [\\ion{S}{ii}] (red), H$\\alpha$ (green), and [\\ion{O}{iii}] (blue), all data from MCELS (see Sect.\\,\\ref{observations_optical}). The red box delineates the area shown in Fig.\\,\\ref{fig_rgb_image}. Noticeable substructures are\\,: DEM L205 (A), the SNR candidate analysed in this paper; N51A (B) and N51C (C, also named DEM L201), two \\ion{H}{ii} regions also seen in the radio and the IR; the SB N51D, or DEM L192, in (D).} \\label{fig_rgb_N51} \\end{figure}   In this paper, we present new X-ray and radio-continuum observations (with \\xmm\\ and ATCA) of the LMC SNR candidate \\dem. Archival optical and infrared (IR) observations are analysed as well. The source lies in a very complex environment, at the eastern side of the \\ion{H}{ii} complex \\mbox{\\object{LHA 120-N 51}} \\citep[in the nebular notation of][] {1956ApJS....2..315H}. It was classified as a ``possible SNR'' by \\citet*{1976MmRAS..81...89D} (from which the identifier ``DEM'' is taken) based on its optical shell-like morphology. In X-rays, the catalogue of \\emph{ROSAT}'s PSPC sources in the LMC \\citep{1999A&AS..139..277H} includes \\object{[HP99] 534}, located within the extent of \\object{DEM L205}. However, the short exposure and large off-axis position prevented any classification of the X-ray source. \\citet{2001ApJS..136..119D} classified \\object{DEM L205} as a superbubble (SB) with an excess of X-ray emission.  In our new and archival observations, we detected the object at all wavelengths. The subsequent analyses allowed us to confirm the SNR nature of the source and estimate some of its parameters. The paper is organised as follows: in Sect.\\,\\ref{observations}, we present our new X-ray and radio-continuum observations, and archival optical and IR data. The X-ray, radio, and IR data analyses are detailed in Sect.\\,\\ref{data}. We discuss the implications of our multi-frequency study in Sect\\,\\ref{discussion}, and we give our conclusions in Sect.\\,\\ref{conclusions}.    \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.49\\hsize]{RGB_paper2_low.ps} \\includegraphics[width=0.49\\hsize]{RGB_paper_mcels_xraycontours_low.ps}  \\includegraphics[width=0.49\\hsize]{RGB_paper2_ratiocontours_low.ps} \\includegraphics[width=0.49\\hsize]{RGB_paper_mips_low.ps} \\end{center} \\sidecaption \\caption{A multicolour view of \\dem. \\emph{Top left}\\,: X-ray colour image of the remnant, combining all EPIC cameras. Data from two overlapping observations are combined and smoothed (see Sect.\\,\\ref{data_xray_image} for details). The red, green, and blue components are soft, medium, and hard X-rays, as defined in the text. The white circle is the 90\\,\\% confidence error of the \\object{[HP99] 534} position and the green cross is the central position of \\object{DEM L205}. The green dashed ellipse (defined in Sect.\\,\\ref{data_xray_image}) encompasses  the X-ray emission and is used to define the nominal centre and extent of the remnant. \\emph{Top right}\\,: The same region of the sky in the light of [\\ion{S}{ii}] (red), H$\\alpha$ (green), and [\\ion{O}{iii}] (blue),where all data are from the MCELS. The soft X-ray contours from the top left image are overlaid. \\emph{Bottom left}\\,: Same EPIC image as above but with [\\ion{S}{ii}]-to-H$\\alpha$ ratio contours from MCELS data. Levels are (inwards) 0.4, 0.45, 0.5, 0.6, and 0.7. \\emph{Bottom right}\\,: The remnant as seen at 24 $\\mu$m by \\emph{Spitzer} MIPS. Optical and IR images are displayed logarithmically.} \\label{fig_rgb_image} \\end{figure*} ",
        "conclusions": "\\label{conclusions} The first observation of our LMC survey with \\xmm\\ included the SNR candidate \\dem\\ in the field of view. In combination with unpublished radio-continuum data and archival optical and IR observations, we have found all classical SNR signatures, namely\\,: \\begin{itemize} \\item extended X-ray emission \\item optical emission with a shell-like morphology and an enhanced [\\ion{S}{ii}]-to-H$\\alpha$ ratio \\item non-thermal and extended radio-continuum emission. \\end{itemize} The source is also detected in the IR where we predominantly observe thermal emission from dust. We can therefore definitely confirm this object as a supernova remnant. A core-collapse supernova origin is favored, in light of the recent burst of star formation and the presence of many massive stars in the close vicinity of the remnant. The SN exploded in a SB, thus expanding in a low density medium. With a size of $\\sim$ 79 \\mbox{$\\times$~64 pc}, \\dem\\ is one of the largest SNR known in the LMC. Given the low plasma temperature ($kT \\sim$ 0.2\\,--\\,0.3 keV), we derived a dynamical age of about 35~kyr. Whilst completing our survey, we can expect to find other similarly evolved remnants, thereby refining the faint end of the size and luminosity distributions of SNRs in the LMC."
    },
    "1208/1208.4181_arXiv.txt": {
        "abstract": "The observed flux of ultra-high energy (UHE) cosmic rays (CRs) guarantees the presence of high-energy cosmogenic neutrinos that are produced via photo-hadronic interactions of CRs propagating through intergalactic space. This flux of neutrinos doesn't share the many uncertainties associated with the environment of the yet unknown CR sources. Cosmogenic neutrinos have nevertheless a strong model dependence associated with the chemical composition, source distribution or evolution and maximal injection energy of UHE CRs. We discuss a lower limit on the cosmogenic neutrino spectrum which depends on the observed UHE CR spectrum and composition and relates directly to experimentally observable and model-independent quantities. We show explicit limits for conservative assumptions about the source evolution. ",
        "introduction": "Cosmogenic neutrinos are produced when UHE CRs interact with the cosmic radiation background while propagating between their sources and Earth. The frequent interactions with the cosmic microwave background (CMB) limits the propagation of nucleons with energies greater than $E_{\\rm GZK}\\simeq40$~EeV to within a few 100~Mpc and is responsible for the so-called Greisen-Zatsepin-Kuzmin (GZK) cutoff of extra-galactic protons~\\cite{Greisen:1966jv,Zatsepin:1966jv}. Mesons produced in these interactions quickly decay and produce an observable flux of cosmogenic (or GZK) neutrinos~\\cite{Berezinsky:1970xj}. In fact, the observed spectrum of CRs extending up to energies of a few 100~EeV shows a suppression above $\\sim E_{\\rm GZK}$ with high statistical significance~\\cite{Abbasi:2007sv,Abraham:2008ru}. This could be an indication that protons are dominating the flux at these energies. In this case the flux of cosmogenic neutrinos is typically large.  However, the experimental situation is less clear. Measurements of the elongation rate distribution of UHE CR showers indicate a transition of their arrival composition from light to heavy within 4-40~EeV~\\cite{Abraham:2010yv,Unger:2011ry}. If a heavy component dominates also at higher energies the prospect for cosmogenic neutrino production is  ``disappointing''~\\cite{Aloisio:2009sj} or at least less favorable than for the proton scenario. A crucial uncertainty of this scenario is the maximal injection energy of the nucleus with mass number $A$; as long as $E_{\\rm max} \\gg AE_{\\rm GZK}$, even this scenario will produce an appreciable amount of cosmogenic neutrinos~\\cite{Ahlers:2011sd}. If this condition is not met interactions with the subdominant cosmic photon background from the optical/infra-red will still contribute to the cosmogenic neutrino flux. We will use the estimate of Ref.~\\cite{Franceschini:2008tp} for our calculation.  The IceCube neutrino observatory has reached the sensitivity for the detection of optimistic cosmogenic neutrino fluxes~\\cite{Abbasi:2011ji}. In the case of a non-observation it is of interest to know a lower limit on the various source emission possibilities for their definite exclusion. Lower cosmogenic neutrino flux limits have already been discussed in the context of proton-dominated scenarios via a deconvolution of early Auger data~\\cite{Fodor:2003ph}. We will discuss in this article updates of these lower limits and extensions to more general assumptions for the source distribution and chemical composition. Similar to Ref.~\\cite{Fodor:2003ph} we will not attempt to construct a specific source emission model that fits the Auger spectrum and elongation rate distribution but we will derive the limits directly from the observed composition measurement and spectrum. From this we can derive a strict lower limit on the cosmogenic flux. ",
        "conclusions": "We have discussed in this article a simple procedure to derive lower limits on the cosmogenic neutrino flux. The limits are based on the observed spectrum and composition of UHE CRs and depend on the unknown evolution of sources. For the case of a proton-dominance in the UHE CR data we show that ARA-37 should identify the flux of cosmogenic neutrinos after 3 years of observation if UHE CR sources follow the star formation rate. For the less optimistic (and less realistic) case of no source evolution it would require 10 years of observation.  In the case of heavy nucleus dominance of the CR flux cosmogenic neutrino predictions are less optimistic. We can derive a lower limit in this scenario by tracking the leading nucleus back to its source. Since photo-disintegration conserves the energy per nucleon of the interaction we can base our analysis on the observed number of nucleons in UHE CRs, which depends on the observed mass composition.  The dominant contribution to the cosmogenic neutrino flux is expected from the proton content in the UHE CR spectrum. We show in Fig.~\\ref{fig2} two cases where we decrease the contribution of protons to 10\\% and 1\\% at 100~EeV and assume source evolution with the star-formation rate. Even this less optimistic case is in reach of ARA-37 after 5 years of observation.  The prediction of cosmogenic neutrinos is very sensitive to the maximal CR injection energy per nucleon. If this is significantly larger than the GZK cutoff, even UHE CR scenarios dominated by heavy nuclei can produce large fluxes of cosmogenic neutrinos. For flat spectra that are sufficiently close to $E^{-2}$ the energy density of these optimistic GZK neutrino predictions depends on the cosmic evolution of the sources.  All cosmogenic neutrino fluxes shown in this analysis are normalized to Auger data. The spectra observed with HiRes and the Telescope are in general larger, which could be a result of an overall systematic energy shift by $20-30$\\%. This corresponds to an upward shift of up to a factor 2 of the energy density $E_{\\rm CR}^2J_{\\rm CR}(E_{\\rm CR})$. Hence the lower limits shown in Figs.~\\ref{fig1} and \\ref{fig2} should be similarly scaled upward.  Finally, we would like to stress that the present analysis does not take into account statistical uncertainties of the CR data. However, the method can be easily extended to this case. In Ref.~\\cite{Ahlers:2010fw} we have shown that an actual fit to HiRes data assuming a proton power-law injection in the sources is statistically consistent with cosmogenic neutrino fluxes that exceed the minimal bound by up to an order of magnitude and are in reach of the IceCube detector.  \\noindent {\\it Acknowledgments.} We thank Albrecht Karle, Ali Kheirandish and Kohta Murase for discussion. MA acknowledges support by a John Bahcall Fellowship for neutrino astronomy of the Wisconsin IceCube Particle Astrophysics Center (WIPAC). FH is supported in part by the National Science Foundation under Grant No.~OPP-0236449, by the DOE under grant DE-FG02-95ER40896 and in part by the University of Wisconsin Alumni Research Foundation."
    },
    "1208/1208.1247_arXiv.txt": {
        "abstract": "Between the BICEP2 and Keck Array experiments, we have deployed over 1500 dual polarized antenna coupled bolometers to map the Cosmic Microwave Background's polarization.  We have been able to rapidly deploy these detectors because they are completely planar with an integrated phased-array antenna.  Through our experience in these experiments, we have learned of several challenges with this technology- specifically the beam synthesis in the antenna- and in this paper we report on how we have modified our designs to mitigate these challenges.  In particular, we discus differential steering errors between the polarization pairs' beam centroids due to microstrip cross talk and gradients of penetration depth in the niobium thin films of our millimeter wave circuits.  We also discuss how we have suppressed side lobe response with a  Gaussian taper of our antenna illumination pattern.  These improvements will be used in Spider, Polar-1, and this season's retrofit of Keck Array. ",
        "introduction": "\\label{sec:intro}  The Cosmic Microwave Background's (CMB) temperature anisotropies and curl-free E-mode polarization anisotropies, both generated by scalar inflationary potentials, have been mapped by numerous experiments and used to constrain a multitude of cosmological parameters.  To date, divergence-free B-modes in the CMB, which would have been seeded by inflationary tensor potentials, have not been detected.  However, current data favor a scalar index $n_s<1$, which suggests that the tensor-to-scalar ratio $r$ may be non-zero \\cite{WMAP_seven_year}.  A B-mode detection at degree angular scales would provide strong confirmation of inflation, but even upper limits on $r$ will constrain the energy scale of inflation, thus making it a fundamental parameter of interest.  Additionally, gravitational lensing can shear E-modes into B-modes, creating a smaller angular scale non-primordial anisotropy.  The details of this peak can constrain the summed mass of the different neutrino species and potentially inform us of the existence of sterile neutrinos\\cite{Dodelson_white_paper}.  Our team has developed a low cost camera design with just enough resolution to detect B-modes from primordial tensor perturbations.  This design builds off the success of BICEP-1, using refracting optics with a 30 arcmin resolution\\cite{Yoon_B1_SPIE}\\cite{Aiken_SPIE}.  BICEP2 and Keck Array's five BICEP2 style cameras are both currently collecting data at the South Pole with nearly 1500 pixels (3000 detectors) collectively \\cite{Keck_array_SPIE}.  The balloon-borne Spider will fly in 2013 from the McMurdo station in Antarctica with six cameras and a comparable pixel count\\cite{Filippini_SPIE_SPIDER}.  Finally, we will deploy the 1.5m crossed-Dragone telescope Polar-1 to the South Pole in 2013 with 1900 pixels.  Its 5 arcmin resolution will be sufficient to detect lensed B-modes.  All four of these experiments use phased-array antenna-coupled TES bolometers developed at Caltech and JPL.  Our detectors are entirely planar, and hence scalable to large monolithic arrays.  This has provided our team with a crucial advantage over competing teams, allowing early deployment of high sensitivity focal planes.  As a result, our technology has matured to a point where we have identified several failure modes from both lab and field data and we have refined our original design with these measurements in mind.  This paper describes some of the efforts to improve our detectors.  The second section summarizes our detector design, the third describes how we corrected differential pointing between polarization pairs, and the last outlines how we have suppressed sidelobe response with a tapered illumination pattern in anticipation of Polar. ",
        "conclusions": ""
    },
    "1208/1208.3132_arXiv.txt": {
        "abstract": "{Using a new catalog of 824 solar and late-type stars with X-ray luminosities and rotation periods we have studied the relationship between rotation and stellar activity. From an unbiased subset of this sample the power law slope of the unsaturated regime, $L_X/L_{bol}\\propto Ro^\\beta$, is fit as $\\beta=-2.70\\pm0.13$. This is inconsistent with the canonical $\\beta=-2$ slope to a confidence of 5$\\sigma$ and argues for an interface-type dynamo. Super-saturation is observed for the fastest rotators in our sample and its parametric dependencies are explored. Significant correlations are found with both the corotation radius and the excess polar updraft, the latter theory being supported by other observations. We also present a new X-ray population synthesis model of the mature stellar component of our Galaxy and use it to reproduce deep observations of a high Galactic latitude field. The model, XStar, can be used to test models of stellar spin-down and dynamo decay, as well as for estimating stellar X-ray contamination rates for non-stellar studies.} ",
        "introduction": "Stars across the Hertzsprung-Russell diagram are known to emit X-rays with only a few exceptions. Solar- and late-type stars are thought to generate their X-rays from a magnetically confined plasma known as a corona (Vaiana et al. 1981), which is believed to be driven by the stellar magnetic dynamo, which itself is thought to be driven by differential rotation in the stellar interior (e.g., Parker 1955), a phenomenon that has been confirmed in the Sun through helioseismology (Duvall et al. 1984). The observed decrease in X-ray emission from the pre-main sequence to the age of the Galactic field population can therefore be attributed to the rotational spin-down of the star, which is driven by mass loss through a magnetized stellar wind (Skumanich 1972).  The relationship between stellar rotation and tracers of magnetic activity is a vital probe of the stellar dynamo. A relationship between rotation and activity was first quantified by Pallavicini et al. (1981), who found that X-ray luminosity scaled as $L_X \\propto (v \\, \\mathrm{sin} \\, i)^{1.9}$, providing the first evidence for the dynamo-induced nature of stellar coronal activity. For very fast rotators the relationship was found to break down with X-ray luminosity reaching a saturation level of $L_X / L_{bol} \\sim 10^{-3}$ (Micela et al. 1985), independent of spectral type. This saturation level is reached at a rotation period that increases toward later spectral types (Pizzolato et al. 2003), but it is unclear what causes this. Despite much work there is yet to be a satisfactory dynamo theory that can explain both the solar dynamo and that of rapidly rotating stars and the continued lack of a sufficiently large and unbiased sample has no doubt contributed to this.  We have produced a new catalog of stars with stellar rotation periods and X-ray luminosities, the details of which are presented in Wright et al. (2011). The catalog includes 824 solar- and late-type stars, including 445 field stars and 379 stars in nearby open clusters (ages 40--700~Myrs). The sample was homogenized by recalculating all X-ray luminosities and converting them all onto the ROSAT $0.1 - 2.4$~keV band. To minimize biases we removed all sources known to be X-ray variable, those that exhibit signs of accretion, or those in close binary systems. The sample is approximately equally distributed across the color range $V-K_s = 1.5 - 5.0$ (G2 to M4) with $\\sim$30 stars per subtype, dropping to $\\sim$10 stars per subtype outside of this, from F7 to M6.  Here we highlight some of the results derived from this sample, focussing attention on the regimes of low-activity dynamo decay and high-activity supersaturation. Details of these findings and further results and discussions from this sample can be found in Wright et al. (2011). Finally we present a new population synthesis model that can be used to address similar problems in stellar activity.  \\begin{figure*} \\begin{center} \\includegraphics[height=450pt, angle=270]{relation.ps} \\caption{$R_X = L_X / L_{bol}$ plotted against $P_{rot}$ (left) and the Rossby number, $Ro = P_{rot} / \\tau$ (right), for all stars in our sample. Known binary stars are shown as plus symbols, and the Sun with a solar symbol. The best-fitting saturated and non-saturated activity--rotation relations are shown as a dashed red line in the right-hand panel.} \\label{allstars} \\end{center} \\end{figure*} ",
        "conclusions": "A new catalog of stars with measured X-ray luminosities and rotation periods is used to study the rotation -- activity relation. The power-law slope of the unsaturated regime is fit as $\\beta = -2.7 \\pm 0.13$, inconsistent with the canonical $\\beta = -2$ value with a confidence of 5$\\sigma$ and arguing for an interface-type dynamo. Coronal supersaturation in ultra-fast rotators is shown to correlate well with both the corotation radius and the excess polar updraft as the parameters of the coronal stripping and convective updrafts theories, though other observations support the latter theory. Finally a new population synthesis model of X-ray emitting stars in a Galactic sightline is introduced and tested successfully on deep observations of the high Galactic latitude COSMOS field. A better fit is found for $\\beta = -2.7$ than for $\\beta = -2$, confirming results with a fixed sample."
    },
    "1208/1208.4524_arXiv.txt": {
        "abstract": "The blazar OJ287 is the most promising (and the only) case  for an extragalactic binary black hole system inspiralling under the action of gravitational radiation reaction. At present, though it is not possible to directly observe the binary components, it is possible to observe the jet emanating form the primary black hole. We argue that the orbital motion of the secondary black hole is reflected in the wobble of the jet and demonstrate that the wobble is orbital position dependent. The erratic wobble of the jet, reported in Agudo et al. (2012), is analyzed by taking into account the binary nature of the system and we find that the erratic component of jet wobble is very small. ",
        "introduction": "OJ287 is a very special case among all quasars: its optical light curve of 120 yr length shows a double periodicity of 60 yrs and 12 yrs [1]. Another remarkable feature of this quasar is that while its radiation is dominated by synchrotron radiation at most times, during the major outbursts its optical to UV spectrum is bremsstrahlung at around $3\\times10^5$ K [2]. These main outbursts occur in unpolarized light, as is expected of bremsstrahlung [3]. Further,  it is the only known quasar showing bremsstrahlung outburst peaks.  Most of these facts were not known in 1995 when a detailed model of this source was constructed [4,5]. The model proposed that the major outbursts which occur at roughly 12 yr intervals are related to a 12 yr period in a black hole binary system. The detailed mechanism for generating the outbursts which occur in pairs separated by 1 to 2 years, was presumed to be impacts on the accretion disk of the primary by the secondary. Hot bubbles of plasma pulled off the accretion disk then generate the bremsstrahlung bursts when the bubbles become optically thin.  At present, the model has been verified in many ways [6] and subsequently, the model is invoked to constrain the spin of the primary black hole and to explore the possibility of testing black hole no-hair theorems [7,8]. It should be noted that the binary black hole model implies typical orbital velocity  $v \\sim 0.1\\,c $ and it explains why higher order post-Newtonian corrections are required to explain observed major outbursts and to make predictions about future major outbursts. In contrast, binary pulsars have $v/c \\sim 10^{-3}$, which is roughly a factor 10 higher than orbital velocity of the Earth. However, it is desirable to probe further observational implications arising from the presence of binary components in the model and this is what is pursued below. ",
        "conclusions": ""
    },
    "1208/1208.5008_arXiv.txt": {
        "abstract": "{We present a photometric catalog of infrared (IR) sources based on the North Ecliptic Pole Wide field (NEP-Wide) survey of AKARI, which is an infrared space telescope launched by Japan. The NEP-Wide survey covered 5.4 deg$^{2}$ area, a nearly circular shape centered on the North Ecliptic Pole, using nine photometric filter-bands from 2 -- 25 $\\mu$m of the Infrared Camera (IRC). Extensive efforts were made to reduce possible false objects due to cosmic ray hits, multiplexer bleeding phenomena around bright sources, and other artifacts. The number of detected sources varied depending on the filter band: with about 109,000 sources being cataloged in the near-IR bands at 2 -- 5 $\\mu$m, about 20,000 sources in the shorter parts of the mid-IR bands between 7 -- 11 $\\mu$m, and about 16,000 sources in the longer parts of the mid-IR bands, with  $\\sim$ 4,000 sources at 24 $\\mu$m. The estimated 5$\\sigma$ detection limits are approximately 21 magnitude (mag) in the 2 -- 5 $\\mu$m bands, 19.5 -- 19 mag in the 7 -- 11 $\\mu$m, and 18.8 -- 18.5 mag in the 15 -- 24 $\\mu$m bands in the AB magnitude scale. The completenesses for those bands were evaluated as a function of magnitude: the 50$\\%$ completeness limits are about 19.8 mag at 3 $\\mu$m, 18.6 mag at 9 $\\mu$m, and 18 mag at 18 $\\mu$m band, respectively. To construct a reliable source catalog, all of the detected sources were examined by matching them with those in other wavelength data, including optical and ground-based near-IR bands. The final band-merged catalog contains about {114,800} sources  detected in the IRC filter bands. The properties of the sources are presented in terms of the distributions in various color-color diagrams. } ",
        "introduction": "An infrared (IR) space telescope AKARI, launched by ISAS/JAXA in 2006 (Murakami et al. 2007), successfully carried out an all-sky survey at mid- and far-infrared wavelengths, as well as several large-area surveys and many other pointed observations across the wavelength range 2 -- 160 $\\mu$m. The North Ecliptic Pole (NEP) survey (Matsuhara et al. 2006) was one of these large area surveys of AKARI. The AKARI telescope had two focal-plane instruments: the Infrared Camera (IRC, Onaka et al. 2007) and the Far-Infrared Surveyor (FIS, Kawada et al. 2007). The FIS covered a 50 -- 200 $\\mu$m range with four wide-band photometric filters, and a Fourier Transform Spectrometer (FTS). The IRC was designed to carry out near- to mid-infrared imaging with nine photometric filters, and spectroscopic observations with a prism and grisms.    \\begin{figure}[h] \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics{fig01}} \\caption{The overall map of the NEP-Wide field. The survey consisted of 446 pointing observations represented by green boxes. Each frame covered a 10$'\\times$10$'$ area with half of its field of view (FoV) overlapped by neighboring frames. The red box and blue lines represent the regions covered by optical surveys at the CFHT and  Maidanak Observatory, respectively. The circular gray-shaded region denotes the NEP-Deep field.} \\label{fig 01} \\end{center} \\end{figure}    The NEP survey is composed of two parts: a wide (henceforth NEP-Wide) survey and a deep (henceforth NEP-Deep) survey. The area observed in the NEP-Wide survey is shown in Fig. 1 (green tiles) and is about 5.4 deg$^{2}$ with a circular shape (whose radius is about 1.25 deg) centered on the NEP ($\\alpha =18^{h}00^{m}00^{s}$, $\\delta= +66^{\\circ}33^{'}38^{''} $), while the NEP-Deep survey covers about 0.6 deg$^{2}$ (Wada et al. 2008, Takagi et al. 2012), with a center slightly offset from the NEP (shaded region).  These surveys were observed using the nine IRC bands to provide nearly continuous coverage from 2 $\\mu$m to 25$\\mu$m.  The filter system is designated as $N2$, $N3$, and $N4$ for the near-infrared (NIR) bands, $S7$, $S9W$, and $S11$ for the shorter part of the mid-IR band (MIR-S), and $L15$, $L18W$, and $L24$ for the longer part of the mid-IR bands (MIR-L) with the numbers representing the approximate effective wavelengths in units of $\\mu$m. The photometric bands with wider spectral widths are indicated by a W at the end. The NEP-Wide survey was carried out with 446 pointed observations, with the field of view (FoV) of each individual frame being 10$'\\times$10$'$. The survey coverage was designed around seven concentric circles and a partial rim in the outermost part.    The details of the strategy, the observational plans for the coordinated pointing surveys, the scientific goals and the technical constraints are described in Matsuhara et al. (2006). The initial results and the catalog for the NEP-Deep survey were reported by Wada et al. (2008). The data characteristics and basic properties of the sources were presented using a subset of NEP-Wide data by Lee et al. (2009). Here, we present the entire data set of NEP-Wide, corrected to remove artificial effects, and provide a more detailed analysis of the photometric results. The main purposes of this paper is to present the data reduction methodology, and provide a point source catalog of the NEP-Wide survey.  This paper is organized as follows. In sect. 2, we present details of the data reduction process, and the additional image corrections needed to improve the efficacy of the image data.  We describe the results of our source extraction and photometry, as well as the properties of the data such as sensitivity and the completeness of  the source detection in \\S 3.  The next section describes the source matching across the available bands to confirm the genuineness of the detected sources in \\S 4. In \\S 5, we present the band-merging procedure and describe the contents in the final catalog. The natures of the detected sources using various color-magnitude diagrams and color-color diagrams are shown in \\S 6.  We summarize  our results in the final section. ",
        "conclusions": "We have carried out the reduction and analysis of the NEP-Wide survey data obtained by the AKARI/IRC. In order to reduce spurious detection, we masked out the regions affected by instrumental effects such as MUXbleeding trails, especially in the NIR bands. The detected sources were compared with the available data at other wavelengths, including optical, ground-based near-IR observations, in addition to the other bands of AKARI. The areal coverage is about a 5.4 deg$^2$ circular field centered on the NEP. The 5$\\sigma$ detection limits of the survey are around 21 AB mag in the NIR bands, 19 -- 19.5 mag in the MIR-S bands, and 18.5 -- 18.8 mag in the MIR-L bands.  The ancillary optical data from the CFHT and the Maidanak observatory are sufficiently deep to identify most of the AKARI sources.  We carried out extensive comparisons by cross-matching of the sources among the photometric bands ranging from optical to mid-IR  wavelengths in order to confirm the validity of the detected sources and exclude the low-reliability sources. Using these results, we produced a band-merged source catalog covering the wavelength bands from the $u^{*}$ band to the MIR $L24$ band. This catalog contains about 114,800 entries. On the basis of this catalog, we have shown the characteristics of the sources using various color-color diagrams. By comparing and using optical stellarity, we found that the NIR and MIR-S band colors provide a reliable means of distinguishing stars from galaxies. Except for the star-like objects, most of the NEP-Wide sources appear to be various types of star forming galaxy. The sources detected in all of the AKARI/IRC bands include interesting sources such as PAH galaxies, AGNs, ULIRGs, DOG candidates, or  MIR-bright early-type galaxies.  The NEP-Wide catalog covers a moderately large sky area with a wide wavelength range. It complements the NEP-Deep catalogs of  Wada et al. (2008) and Takagi et al. (2012), which have better sensitivity and smaller angular coverage, but the same filter bands.  The Spitzer space telescope has also carried out large area surveys such as the Spitzer Wide-Area Infrared Extragalactic (SWIRE) Survey and First Look Survey (FLS). These surveys are carried out with all wide-band filters of Spitzer: 3.6 4.5, 5.8, and 8.0 $\\mu$m with the Infrared Array Camera (IRAC) and 24, 70, and 160 $\\mu$m with the Multiband Imaging Photometer for Spitzer (MIPS). This should be compared to the nearly continuous wavelength coverage of AKARI's NEP surveys from 2.4 to 24 $\\mu$m. The FLS covered the area of about 5 deg$^2$, which is similar to that of the NEP-Wide. The survey area of SWIRE is about ten times larger, but composed of several different fields.  This work is based on observations with AKARI, a JAXA project with the participation of ESA, universities and companies in Japan, Korea, the UK, and Netherlands. This work contains many data obtained by the ground-based  Maidanak Observatory's 1.5 m, KPNO 2.1 m, and CFHT 3.5 m telescopes. This work was supported by the Korean Research Foundation grant 2006-341-C00018. AS and AP have been supported by the research grant of the Polish National Science Centre N N203 51 29 38."
    },
    "1208/1208.0351_arXiv.txt": {
        "abstract": "The ambient hot intra-halo gas in clusters of galaxies is constantly fed and stirred by in-falling galaxies, a process that can be studied in detail with cosmological hydrodynamical simulations. However, different numerical methods yield discrepant predictions for crucial hydrodynamical processes, leading for example to different entropy profiles in clusters of galaxies. In particular, the widely used Lagrangian smoothed particle hydrodynamics (SPH) scheme is suspected to strongly damp fluid instabilities and turbulence, which are both crucial to establish the thermodynamic structure of clusters. In this study, we test to which extent our recently developed Voronoi particle hydrodynamics (VPH) scheme yields different results for the stripping of gas out of infalling galaxies, and for the bulk gas properties of cluster. We consider both the evolution of isolated galaxy models that are exposed to a stream of intra cluster medium or are dropped into cluster models, as well as non-radiative cosmological simulations of cluster formation. We also compare our particle-based method with results obtained with a fundamentally different discretisation approach as implemented in the moving-mesh code {\\small AREPO}. We find that VPH leads to noticeably faster stripping of gas out of galaxies than SPH, in better agreement with the mesh-code than with SPH. We show that despite the fact that VPH in its present form is not as accurate as the moving mesh code in our investigated cases, its improved accuracy of gradient estimates makes VPH an attractive alternative to SPH. ",
        "introduction": "In the hierarchical structure formation model \\citep{White1991}, interactions or mergers of galaxies with other galaxies are common processes, and depending on the environment, a diverse set of physical processes can become important.  For example, when a galaxy falls into a cluster of galaxies, gravity may affect its integrity through gravitational tidal forces. The closer the galaxy gets to the cluster the more it is exposed to the hot intra-cluster medium (ICM). This is experienced as a headwind by the galaxy that compresses and removes the gas of the galaxy \\citep{GunnGott}. This may lead to the formation of a bow shock in front of the galaxy if it supersonically ploughs through the ICM.  Whereas the gas at the front of the galaxy is primarily compressed by ram pressure in this situation, the sides are exposed to strong shear flows which may trigger fluid instabilities, leading to stripping and mixing of the gas with the ICM. This in turn crucially determines how rapidly star formation is turned off in an infalling galaxy, and how metals are mixed into the group or cluster ICM \\citep[see e.g.][]{Larson1980, Balogh2000, Domainko2006}. The hydrodynamical interaction processes between galaxies and groups/clusters may thus influence important observational effects, such as the density morphology relation \\citep{Dressler1980, vanderWel2010, Boselli2006}, or the existence of a tight red cluster galaxy sequence \\citep{Baldry2004, Cortese2009}.  Another interesting aspect besides the stripping is the generation of turbulence by galaxies infalling into clusters.  Numerical simulations \\citep{Dolag2005, Vazza2006, Iapichino2008, Dolag2009} indicate that turbulence can be generated in the wake of an infalling galaxy.  This interesting phenomenon may play an important role in determining the dynamics of the cluster gas and could in principle be detected through X-ray observations of the broadening of sharp metal lines \\citep{Sunyaev2003}.  Modelling these processes reliably goes back to \\citet{GunnGott} and is a considerable challenge, as direct hydrodynamic simulations of the relevant processes are very difficult to model due to the large dynamic range, and the uncertainties associated with the modelling of star formation and feedback processes.  Furthermore, it is not clear whether the hydrodynamical techniques presently in use are strongly affected by systematic errors in this regime \\citep [e.g.][]{Agertz}.  Indeed, previous studies of galaxy-cluster interactions with different numerical schemes disagreed strongly about how fast a galaxy loses its gas to the cluster. For example, a grid-based simulation of \\citet{Quilis2000} predicted a much higher stripping rate than SPH-based work by \\citet{Abadi1999}.  In this paper, we concentrate on the numerical aspects of the stripping of a galaxy's gas as it interacts with the ICM. \\citet{Agertz} has recently shown that the popular smoothed particle hydrodynamics (SPH) technique has problems to properly account for fluid instabilities, prompting significant concerns about a possible unphysical suppression of stripping processes \\citep[see also][]{Price2008, Wadsley2008,Read2010,Cha2010,Valcke2010,Junk2010,Abel2011}. In a recent study \\citep{Hess2010} we have therefore proposed a new `Voronoi particle hydrodynamics' (VPH) method that improves on the widely used SPH technique \\citep{Lucy1977, Gingold_Monaghan, Larson1978} in several respects. By employing a Voronoi tessellation for the local density estimate, a consistent decomposition of the simulation volume is achieved, contact discontinuities can be resolved much more sharply, and a `surface tension' effect across them is avoided. Our preliminary tests of VPH based on the `blob-test' of \\citet{Agertz} already suggested that stripping is more efficient in VPH compared with SPH. Here we shall investigate this in more detail using more realistic set-ups that mimic galaxy evolution processes. In addition to the two particle-based Lagrangian schemes for hydrodynamics, SPH and VPH, we will also carry out comparison simulations with the moving-mesh code {\\small AREPO} \\citep{Springel2010}. Whereas {\\small AREPO} uses a Voronoi tessellation as well, this code employs an entirely different methodology for fluid dynamics, based on a finite volume Godunov scheme that calculates hydrodynamical fluxes with a Riemann solver across mesh boundaries. The comparison of this diverse set of three numerical methods is useful to understand and quantify the systematic uncertainties of the different methods.  In our tests simulations, we ideally want to investigate as realistic conditions as possible, including a full treatment of gravity and dark matter, as well as star formation and cooling. This in principle calls for simulations in which a well-resolved galaxy model is dropped into an equally well represented cluster of galaxies. Because the substantial computational cost of such a set-up severely limits the resolution that can be achieved, we base part of our study on a number of more idealised simulations, for example by placing galaxy models into a `wind tunnel' that mimics the impinging flow of gas onto a galaxy in orbit in a cluster. Finally, we also carry out cosmological simulations of the formation of the `Santa Barbara cluster' \\citep{Frenk}, primarily to study how well VPH performs in non-radiative cosmological simulations of cluster formation compared to the other techniques.  Already in past studies, the Santa Barbara cluster comparison project has led to important insights into how numerical effects impact the thermodynamic structure of simulated clusters, and for this test problem, there is already a considerable body of results in the literature \\citep[e.g.][]{Wadsley2004,gadget2,Thacker2006}.  This paper is structured as follows. In Section~\\ref{SecBasics}, we introduce the numerical methods we will use in our numerical comparison study. In Section~\\ref{IsolatedGal}, we check how well the different numerical schemes represent the evolution of an isolated galaxy, and whether there are already significant differences at this level. We then consider in Section~\\ref{WindTunnel} wind-tunnel experiments in which we expose the isolated galaxy models to a supersonic wind, allowing a detailed examination of the stripping process. In Section~\\ref{Infall}, we follow up on these experiments by studying the behaviour of a galaxy model that falls into an isolated spherical cluster, comparing the VPH results with those obtained with SPH and {\\small AREPO}.  Finally, in Section~\\ref{Cosmo}, we investigate the performance of VPH in cosmological simulations of cluster formation, using re-simulations of rich clusters identified in the Millennium Simulation as well as the Santa Barbara cluster initial conditions. We give a discussion of our results in the context of a vortex test problem in Section~\\ref{SecConclusions}, combined with a summary of our conclusions. ",
        "conclusions": "\\label{SecConclusions}   In this study, we have carried out a systematic comparison of the properties of our new Voronoi particle hydrodynamics (VPH) method \\citep{Hess2010} with respect to standard SPH and moving-mesh hydrodynamics as implemented in the {\\small AREPO} code. We have focused on stripping processes in galaxies and halos upon in-fall into galaxy clusters. Here, it is expected that the outer parts of gaseous disks are quickly removed due to ram pressure stripping \\citep{GunnGott1972}, but the subsequent more gradual gas loss sensitively depends on the ability of hydrodynamical codes to capture fluid instabilities occurring in shear flows around the galaxies. The recent findings that SPH appears to exhibit severe inaccuracies in this regime has prompted us to develop the alternative VPH method. In the present paper, we have studied how this new technique compares with traditional SPH and the new moving-mesh code technique in a number of basic problems relevant for galaxy formation.   \\begin{figure} \\begin{center} \\resizebox{8.5cm}{!}{\\includegraphics{plots/gresho_profile_sph.eps}}\\\\ \\resizebox{8.5cm}{!}{\\includegraphics{plots/gresho_profile_vph.eps}}\\\\ \\resizebox{8.5cm}{!}{\\includegraphics{plots/gresho_profile_arepo.eps}} \\caption{Radial profiles of the azimuthal velocity in evolved simulations ($t=1.0$) of the Gresho vortex test, calculated at resolution $80\\times 80$ with different numerical techniques (SPH, VPH, and {\\small AREPO}), as labelled. The small black dotes show individual velocities of simulation particles/cells, while the red dots give the averaged solution binned in radial annuli.  The blue thick lines give the analytic solution that is realised in the initial conditions.} \\label{Gresho_profiles} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\resizebox{8.5cm}{!}{\\includegraphics{plots/gresho_L1norm.eps}} \\caption{Convergence rate of three different numerical techniques for the Gresho vortex test. The black points refer to measurements of the L1 error norm at $t=2.0$ for the SPH method, with resolutions ranging from $20\\times 20$ to $640\\times 640$. The L1 error has been calculated on the basis of all single particles/cells. The blue dots give the corresponding measurements for our VPH method, while the red points are for the {\\small AREPO} moving-mesh code. The dashed lines indicate power-laws with slopes $-0.3$, $-0.7$ and $-1.4$ for SPH, VPH and {\\small AREPO}, respectively.} \\label{GreshoConvRate} \\end{center} \\end{figure}    To this end, we first compared results for isolated compound galaxy models, both in isolation and in wind tunnels where they were exposed to a supersonic head wind. This set-up allowed relatively high-resolution simulations of wind--ISM interactions. Our simulations have revealed non-negligible differences in the rate at which dense ISM gas is stripped and dispersed, and in the appearance of this gas in the downstream part of the flow. SPH showed a lower stripping rate than both VPH and the mesh-code {\\small AREPO}. As a result, the SPH galaxy also experienced the largest displacement due to the ram pressure of the impinging wind. Also, we were able to show that essentially none of the ISM gas in SPH and VPH could ever be transferred to much lower density. Instead, if gas was stripped, it stayed in coherent dense blobs, where even star formation could continue. Furthermore due to elevated pressure caused by the surface tension effect in SPH, the star formation in these stripped blobs remained at an unphysically high level. {\\small AREPO}, in contrast, showed a rapid loss of gas out of the disk, which was furthermore efficiently mixed with other gas, so that lower densities were reached quickly by the stripped gas and star formation was stopped.  We followed up these simulations with numerical experiments where we dropped galaxy models in live cluster models. Even though here the resolution was substantially lower, we obtained results in good qualitative agreement with our wind tunnel runs. Likewise, in non-radiative cosmological simulations of galaxy cluster formation, we followed individual subhalos as they fell into the forming cluster, finding again that the gas content of satellite systems declined slowest in SPH, while the stripping in VPH and especially in {\\small AREPO} proceeded noticeably faster.  Finally, we considered simulations of the Santa Barbara cluster, which has become an important test problem for evaluating cosmological hydrodynamical codes. While the VPH runs revealed a slightly elevated entropy compared to SPH at the smallest radii, they in general agreed quite well with SPH and in particular did not provide evidence at the highest resolution that an entropy core similar to those found in mesh-codes such as {\\small AREPO} is formed. This is perhaps to be expected if the entropy core indeed primarily arises from mixing processes \\citep{Mitchell2009} that are largely absent in SPH and VPH, by construction.  Overall, it thus appears that VPH offers some improvements over SPH without however changing its fundamental character. We argue that the most important origin of these differences lies in an improved gradient estimate in VPH compared to SPH. VPH is second-order accurate in the gradient estimates, i.e.~a linear gradient is always reproduced exactly, {\\em independent} of the particle distribution \\citep{Springel2010}. In contrast, SPH has a so-called zero-th order error in its gradient estimate \\citep[e.g.][]{Read2010}. This in particular means that even for equal pressures for all particles the pressure force not necessarily vanishes \\citep{Abel2011}, and furthermore, the absolute size of the gradient error grows with the pressure itself. The gradient errors in SPH have also been linked to the creation of small-scale velocity noise in studies of subsonic turbulence \\citep{Bauer2011}.  \\label{gresho_in_conclusion}  A good test problem for appreciating this difference in the gradient accuracy is the flow of a stable vortex as suggested by \\citet{Gresho1990}.  In this `triangular vortex problem' a fluid is set up with an azimuthal velocity profile (see Appendix~\\ref{gresho_appendix} for details), such that the pressure gradients counter the centrifugal force, and the vortex evolves in a time-independent, stable fashion.  Figure~\\ref{Gresho_profiles} shows the radial velocity profile after the vortex has been evolved for a time $t=1$ with the VPH, SPH and {\\small AREPO} codes in 2D, using a $80\\times 80$ Cartesian grid for the initial conditions in the domain $[-0.5,0.5]^2$.  It can be seen clearly that SPH shows a much larger velocity scatter than the other two codes, and its solution has already degraded quite noticeably relative to VPH and {\\small AREPO}. Especially in the inner domain, where the fluid rotates like a solid body, SPH deviates systematically from the analytic solution. We note that part of this degradation can be influenced by the artificial viscosity setting \\citep{Springel2010b,Read2011}, but if a higher viscosity is used to suppress the velocity noise it typically also leads to a faster viscous erosion of the vortex. In VPH, some velocity scatter is seen as well, but it is appreciably smaller than in SPH, which we interpret as a consequence of the more accurate gradient estimates in VPH.  Arguably one of the best ways to quantify the accuracy of the different numerical techniques for this analytic test problem is to look at an objective error measure as a function of resolution. In Figure~\\ref{GreshoConvRate}, we compare the L1-error norm for the azimuthal streaming velocity of the Gresho test as a function of resolution for all three techniques. It is evident that ordinary SPH converges only very slowly, whereas VPH shows a considerably improved convergence rate. This directly demonstrates an important improvement brought about by VPH, which we can directly trace to better gradient estimates. The latter appears also as the primary reason for the better results for stripping we obtained with VPH compared to SPH, which is due to the more accurate treatment of contact discontinuities and the avoidance of the `gap' phenomenon of SPH.  Note that the difference in convergence rate also means that VPH will outperform SPH in terms of computational cost if very high accuracy needs to be achieved, provided its computational cost differs only by a constant factor of order unity which is indeed the case in our present implementation. However, the Voronoi mesh construction adds substantial computational cost compared to ordinary ``standard'' SPH, making the VPH code about a factor 3-4 slower for pure hydrodynamics when the same number of particles is used. This is mitigated to some extent if self-gravity is included (which is typically about as expensive or slightly more expensive than SPH-based hydrodynamics), reducing the difference to less than a factor of 2. We note that some alternative suggestions to improve standard SPH, such as the scheme by \\citet{Read2010} which involves a dramatic increase of the number of smoothing neighbours, also require an increased computational cost per resolution element. Which of these schemes is ultimately the most efficient one (in the sense of reaching a given accuracy goal with the smallest computational cost) is difficult to answer in general, and is in any case a problem-dependent and implementation-dependent question.   According to Fig.~\\ref{GreshoConvRate}, VPH still falls short of the better convergence rate obtained with the moving-mesh code {\\small AREPO} \\citep[and similarly with fixed grid mesh codes such as {\\small ATHENA}, see][]{Stone2008,Springel2010}.  The above discussion suggests that higher order density estimates combined with at least equally accurate gradient estimates are needed to improve on SPH and VPH in this respect. Some suggestions in this direction have recently been made \\citep[e.g.][]{Read2011,Maron2011,McNally2011}, and it will be interesting to see whether they can successfully yield significant accuracy improvements in cosmological applications such as those discussed here."
    },
    "1208/1208.5929_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sectintro} The European Space Agency (ESA) Gaia mission, scheduled for launch in Spring 2013, is a space-based all-sky survey. The Gaia spacecraft will provide astrometry, photometry and spectroscopy for point-like sources down to $V\\sim20$. Gaia's science data comprises absolute astrometry, broad-band photometry, and low-resolution spectro-photometry. During its 5-year mission, Gaia will survey about 1 billion stars and 300,000 Solar System objects, of which the majority will be Main Belt asteroids. It will also survey about 500,000 point-like extragalactic sources and $\\sim1$ million faint galaxies. The astrometric precision for the mission will be better than $10\\mu as$ for stars brighter than $V\\sim13$ and about $25\\mu as$ for stars $V\\sim15$. Gaia will initially build on, and then surpass, the results of the Hipparcos mission of about 20 years ago \\citep{Mignard11a}. Such an observational effort has been compared to mapping the human genome, for the amount of collected data and for the impact that it will have on all branches of astronomy and astrophysics.  In addition to the above core science goal, the Gaia space mission will discover exotic transient objects in large numbers. Many thousands of transients will be discovered including exoplanets and supernovae. Tens of thousands of brown and white dwarfs will be identified spectroscopically and, within our Solar System, some hundreds of thousands of minor planets will be observed. Of particular interest will be the numbers of unusual minor planets such as minor planets which have high inclinations, so that they are normally outside the regions of sky routinely surveyed by Near-Earth Asteroid (NEA) programmes, and inner Solar System Trojan asteroids.  Because Gaia's primary mission is to perform a space-based all-sky survey, it is not designed to conduct any targeted follow-up studies. Gaia will be constrained by its orbit and by its design to survey the sky as completely as possible \\citep{2007EM&P..101...97M}. For these reasons, and because of the expected vast number of discoveries of transient phenomena, Gaia will not be able to either confirm or perform detailed studies of the discoveries.  \\subsection{Gaia Follow-Up Network}  During its 5 year mission, Gaia will observe many new SSOs (Solar System Objects). One feature of Gaia observations is the ability to image at rather low solar elongation (45 degrees), enabling detections of Earth-crossing asteroids (Atens) and Inner Earth Asteroids (IEAs), and discoveries of new NEAs at larger solar elongation. In performing an all-sky survey Gaia will necessarily survey regions of sky away from the regions targeted by NEA programmes. This holds the potential to discover asteroids with high orbital inclinations, and Trojan asteroids in the stable Lagrangian regions in the orbits of the planets in the inner Solar System \\citep{2011arXiv1111.2427T,2012MNRAS.420L..28T,2012MNRAS.424..372T}.  In such cases, due to the motion of the objects and limiting magnitude, the scanning law of Gaia will restrict the orbit determination to be constrained on a very small number of astrometric measurements. This implies that only a ground-based optical network can obtain accurate orbital modeling, based on enough astrometric measurements. This is the primary objectives of the Gaia Follow-Up Network for Solar System Objects (FUN-SSO).  In addition to the improvement in accuracy of the astrometric data used for orbital modeling of specific objects, some rare and peculiar SSOs such as asteroids with cometary activity could be studied \\citep{2008P&SS...56.1812T}. Due to the limitations of Gaia's observing method, a ground-based follow-up network will be crucial for studying the physical characteristics of these objects.  In order to be effective in acquiring high-quality follow-up imaging following an alert by the Gaia data processing system, the network must have a large geographical coverage. This is why several observing stations have been invited for participation in this project. In 2010, the Gaia DPAC (Data Processing and Analysis Consortium) initiated a program to indentify optical telescopes for a Gaia follow-up network for SSOs to perform the following critical tasks: confirmation of alert, identification of transient, continuous monitoring and tracking. At the Gaia Follow-up Network Meeting in Paris 2010, the Zadko Telescope was highlighted as an important contribution to this network, because of its longitude and robotic operation. To date the network comprises 37 observing sites (representing 53 instruments) (cf. figure 1).  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.80]{Gaia-FUN-SSO-May2012.eps} \\caption{ \\small The global distribution of observing stations participating in the Gaia Follow-Up Network of Solar System Objects. The distribution is biased towards northern latitudes and European longitudes. Blue markers are the current observing stations, red markers are existing facilities which have expressed interest but are not confirmed, and yellow markers are planned future facilities.} \\end{center} \\end{figure*}  SSOs and stars are not differentiated on board but are treated the same way. Each source brighter than $V = 20$ is identified by the star mapper (SM) on the CCD array. A window is defined for each identified source for accumulation of charge as it crosses the CCD. An automated data reduction pipeline (implemented by the science consortium, DPAC) is run over all sources. Those which are found to move (from one epoch to the other, essentially) or are extended are sent to a specific branch of the pipeline which applies a specialized data reduction including identification and processing of SSOs. Most of the SSOs that Gaia will observe will be known, so they will be simply identified by comparison to ephemeris derived from the complete Minor Planet Center (MPC) catalogue. The remaining objects which are moving and not identified are classed as new asteroid candidates. For these, a tentative orbit is computed for ground-based recovery and an alert will be issued to the Gaia-FUN-SSO network.   Alerts from the Gaia data processing system will be received between 24 and 48 hours after detection.  Problems may arise for observations of peculiar objects: fast moving objects, faint objects, NEAs close to Earth with strong parallaxes. Hence, small and less sensitive telescopes ($<0.6$ m) can be useful but more sensitive instruments ($>1$ m) will be essential. Furthermore, fully robotic telescopes are proven as ideal for observations triggered by an alert. Among the telescopes in the Gaia-FUN-SSO network, five are robotic telescopes\\footnote{\\url{http://www.imcce.fr/gaia-fun-sso/} for an updated list}. ",
        "conclusions": "During its 5 year mission, Gaia will discover new SSOs, including inner earth asteroids (IEAs) and Earth-crossers (Atens), and new NEAs at larger solar elongations. Because the scanning law of Gaia will restrict the orbit determination to a very small number of astrometric measurements, only a ground-based optical network can obtain accurate orbital modeling, based on enough astrometric measurements. This is the primary objective of the Gaia Follow-Up Network for Solar System Objects (FUN-SSO).  In 2010, the Gaia DPAC (Data Processing and Analysis Consortium) initiated a program to identify ground-based optical telescopes for the Gaia FUN-SSO. The main criteria for the network are facilities that are ideally robotic, or can be re-scheduled to give priority to a Gaia alert. The Zadko Telescope was identified as a potentially important contributor to the network because of its robotic operation, sensitivity and location. It will used for validation of new SSOs, astrometry and photometry of these targets.  In 2011, the Gaia FUN-SSO working group initiated an observing campaign of SSOs, with the main goal to test the capability of the network node facilities. Simulated Gaia alerts were distributed to the network (initially by email) providing ephemerides for follow-up imaging of real targets. We tested the Zadko Telescope capability for participating in the observing campaign of PHA 2005 YU55 in late 2011. Of the 37 observing stations in the network, 11 were able to participate and provide useful astrometric data. The Zadko Telescope contributed the third-longest observational coverage from 14 participating telescopes, and was the only contribution from a Southern Hemisphere telescope. The data from the observations showed that the Zadko Telescope scheduling system worked successfully, with accurate telescope pointings, and positional data that was used for astrometry with an uncertainty of about 0.2 arcseconds. Testing of the networks ability to constrain orbits of SSOs will continue until the launch of Gaia.  Following upgrades to the Zadko Telescope observatory in 2012-2013, the telescope will resume operation several optical transient projects. These improvements will also provide scope for participation in additional projects, including follow-up of exo-planet candidates. Furthermore, an image processing pipeline for automated identification and analysis of optical transients is in development and planned for implementation in 2013."
    },
    "1208/1208.0810_arXiv.txt": {
        "abstract": "Fermi-LAT spectra at high energies (HE, 0.1-100 GeV) are often extrapolated to very high energies (VHE, $\\gtrsim$100 GeV),  and considered either a good estimate or an upper limit for the blazars intrinsic VHE spectrum.  This assumption seems not well justified, neither theoretically nor observationally. Besides being often softer, observations do indicate that spectra at VHE could be also harder than at HE, even when adopting the limit of $\\Gamma\\geq1.5$. Results based on such straightforward GeV-TeV extrapolations are in general not reliable, and should be considered with caution. ",
        "introduction": "The well-determined Fermi-LAT spectra (or Upper Limits, UL) in the MeV-GeV band for several TeV blazars are often used to derive stronger constraints on the diffuse  extragalactic background light (EBL), or to constrain the distance for BL Lacs of uncertain redshift (e.g. \\cite{lat1553,ver1424,persic,orr,prandini}).  The underlying assumption of these studies is that the extrapolation of the Fermi-LAT spectrum to the VHE band is either a good estimate or an upper limit for the intrinsic VHE spectrum of the source, since they belong to the same hump in the spectral energy distribution (SED).  The extrapolation is done either as a simple extension of the power-law (or best fit) model for the Fermi-LAT data, or using a one-zone synchrotron self-Compton model (SSC) on the overall SED.  However, the observational evidence shows that this assumption is not well justified anymore. ",
        "conclusions": ""
    },
    "1208/1208.1816_arXiv.txt": {
        "abstract": "% {\\cthree\\ is the smallest pure carbon chain detected in the dense environment of star forming regions, although diatomic C$_2$ is detected in diffuse clouds.  Measurement of the abundance of \\cthree\\ and the chemistry of its formation  in dense star forming regions has remained relatively unexplored.} {We aim to identify the primary \\cthree\\ formation routes in dense star forming regions  following a chemical network producing species like CCH and \\cthreehtwo\\ in the star forming cores associated with DR21(OH), a high mass star forming region.  } {We have observed velocity resolved spectra of four ro-vibrational far-infrared transitions of \\cthree\\ between the vibrational ground state and the low-energy $\\nu_2$ bending mode  at frequencies between 1654--1897\\,GHz using HIFI on board {\\em Herschel}, in DR21(OH). Several transitions of CCH and \\cthreehtwo\\ have also been observed with HIFI and the IRAM 30m telescope. Rotational temperatures and column densities for all chemical species were estimated.  A gas and grain warm-up model was used to obtain estimates of  densities and temperatures of the envelope.  The chemical network in the model has been used to identify the primary \\cthree\\ forming reactions in DR21(OH).} {We have detected \\cthree\\ in absorption in four far-infrared transitions, $P(4)$, $P(10)$, $Q(2)$ and $Q(4)$. The continuum sources MM1 and MM2 in DR21(OH) though spatially unresolved, are sufficiently separated in velocity to be identified in the \\cthree\\ spectra. All \\cthree\\ transitions are detected from the embedded source MM2 and the surrounding envelope, whereas only $Q(4)$ \\& $P(4)$ are detected toward the hot core MM1. The abundance of \\cthree\\ in the envelope and MM2 is $\\sim 6\\times10^{-10}$ and $\\sim 3\\times10^{-9}$ respectively.  For CCH and \\cthreehtwo\\ we only detect emission from the envelope and MM1.  The observed CCH, \\cthree\\, and \\cthreehtwo\\  abundances are most consistent with a chemical model with $n_{\\rm H_2}\\sim$ 5$\\times 10^{6}$~\\cmcub, post-warm-up dust temperature, $T_{\\rm max}$=30~K and a time of $\\sim$ 0.7--3~Myr.}  {Post warm-up gas phase chemistry of CH$_4$ released from the grain at $t \\sim$ 0.2\\,Myr and lasting for 1\\,Myr can explain the observed \\cthree\\ abundance in the envelope of DR21(OH) and no mechanism involving photodestruction of PAH molecules is required. The chemistry in the envelope is similar to the warm carbon chain chemistry (WCCC) found in lukewarm corinos.  We interpret the observed lower \\cthree\\ abundance in MM1 as compared to MM2 and the envelope to be due to the destruction of \\cthree\\ in the more evolved MM1.  The timescale for the chemistry derived for the envelope is consistent with the dynamical timescale of 2\\,Myr derived for DR21(OH) in other studies.  } ",
        "introduction": "\\begin{table*}[t] \\begin{center} \\caption{Properties of the sources in DR21(OH). \\label{tab_srcprop}} \\begin{tabular}{cccrrrrrc} \\hline \\hline Component & $\\alpha$(2000) & $\\beta$(2000) & $v_{\\rm LSR}$ & Mass$^{a}$ & $T_{\\rm d}$ & N(H$_2$)$^{b}$ & $<n(H_2)>$ & References\\\\ & & & (km~s$^{-1}$) & (\\msun) & & (10$^{23}$\\cmsq) & (10$^7$ \\cmcub) &\\\\ \\hline MM1 & 20$^{\\rm h}$39$^{\\rm m}$01\\fs0 & 42\\arcdeg22\\arcmin48\\arcsec & $-4.1\\pm$0.3 & 88 & 58 & 4.5 & 4.1 & 1,2\\\\ MM2 & 20$^{\\rm h}$39$^{\\rm m}$00\\fs4 & 42\\arcdeg22\\arcmin43\\farcs8 & $-0.7\\pm$0.3 & 143 & 30  & 7.4 & 6.8 & 1,2\\\\ Envelope & \\ldots & \\ldots & $-3.1\\pm$0.3 & 118$^{c}$  & & 3.5 & 0.3 & 1,3,4\\\\ \\hline \\end{tabular} \\end{center}  $^{a}$ All estimated for a distance of 1.5~kpc\\\\ $^{b}$ Within a 15\\arcsec\\ beam\\\\ $^{c}$ Mass of the envelope derived by removing the contributions of MM1 and MM2 from the total mass of the region estimated by \\citet{motte2007}. \\\\ References: (1) \\citet{mangum1991}, (2) \\citet{mangum1992}, (3) \\citet{wilson1990}, (4)\\citet{motte2007}. \\end{table*}   \\begin{table*}[t] \\begin{center} \\caption{Spectroscopic parameters for the \\cthree\\ \\& CCH transitions observed with HIFI/Herschel \\label{tab_specdata}} \\begin{tabular}{lcllrc} \\hline \\hline Species & Transition   & Frequency & {A-coeff} & E$_l$  & Beam Size\\\\ &              & (MHz)     & (s$^{-1}$)  &  (cm$^{-1}$) & \\\\ \\hline \\cthree, ($J$,$v$) & (9,1) -- (10,0) P(10) & 1654081.66 & 2.38~10$^{-3}$  & 47.3 & 13\\arcsec\\\\\\ & (3,1) --  (4,0) P(4) & 1787890.57 & 2.72~10$^{-3}$  & 8.6 & 12\\arcsec\\\\\\ & (2,1) --  (2,0) Q(2) & 1890558.06 & 7.51~10$^{-3}$  & 2.6 & 11\\arcsec\\\\\\ & (4,1) --  (4,0) Q(4) & 1896706.56 & 7.58~10$^{-3}$  & 8.6& 11\\arcsec\\  \\\\ \\hline CCH, N$_{J,F}$  & 6$_{13/2,7}$ -- 5$_{11/2,6}$  & 523971.5704 & 4.58~10$^{-4}$ & 43 & 40\\arcsec \\\\ & 6$_{13/2,6}$ -- 5$_{11/2,5}$  & 523972.1630 & 4.53~10$^{-4}$ & 43 & 40\\arcsec \\\\ CCH, N$_{J,F}$  & 6$_{11/2,6}$ -- 5$_{11/2,5}$  & 524033.9075 & 4.51~10$^{-4}$ & 43 & 40\\arcsec\\\\ & 6$_{11/2,5}$ -- 5$_{9/2,4}$   & 524034.5305 & 4.43~10$^{-4}$ & 43 & 40\\arcsec\\\\  \\hline \\end{tabular} \\end{center} $^\\ast$ Herschel HIFI beamsizes are taken from \\citet{roelfsema2012}. \\end{table*} \\begin{table*}[t] \\begin{center} \\caption{Spectroscopic \\& Observational parameters for the transitions of CCH and \\cthreehtwo\\  observed with IRAM 30~m. \\label{tab_iramspec}} \\begin{tabular}{lcllrccc} \\hline \\hline Species & Transition   & Frequency & {A-coeff} & E$_l$ & $\\theta_{\\rm FWHM}$ & F$_{\\rm eff}^a$ & B$_{\\rm eff}^b$\\\\ &              & (MHz)     & (s$^{-1}$)  &  (cm$^{-1}$) & (\\arcsec) & & \\\\ \\hline CCH, N$_{J,F}$ & 1$_{3/2,1}$ -- 0$_{1/2,1}$ &  87284.1050 & 2.60~10$^{-7}$ & 0.0015& 29\\arcsec & 0.95 & 0.75\\\\ & 1$_{3/2,2}$ -- 0$_{1/2,1}$ &  87316.8980 & 1.53~10$^{-6}$ & 0.0015& 29\\arcsec & 0.95 & 0.75\\\\ & 1$_{3/2,1}$ -- 0$_{1/2,0}$ &  87328.5850 & 1.27~10$^{-6}$ & 0.0015& 29\\arcsec & 0.95 & 0.75\\\\ \\hline \\cthreehtwo, ($J_{K_a},{K_c}$) & 2$_{1,2}$ -- 1$_{0,1}$ & 85338.8930 & 2.32~10$^{-5}$ & 1.7 & 29\\arcsec & 0.95 & 0.75\\\\ \\hline \\hline \\end{tabular} \\end{center} $^a$ $F_{\\rm eff}$ Forward efficiency and $^b$ $B_{\\rm eff}$  Beam efficiency of the telescope. \\end{table*}  Small carbon chains are important for the chemistry of stellar and interstellar environments for several reasons: they are ubiquitous throughout the  interstellar medium \\citep{adamkovics2003}, they  are likely to participate in the formation of long carbon chain molecules, and they are products of photo-fragmentation cascades of polycyclic aromatic hydrocarbons (PAHs) \\citep{radi1988,pety2005}.  Triatomic carbon, \\cthree, was first tentatively identified in interstellar gas by \\citet{vanorden1995} and \\citet{haffner1995}.  The mid-infrared spectrum of \\cthree\\ ($\\nu_3$ antisymmetric stretching mode) was measured in the circumstellar envelope of CW Leo (IRC +10216) by \\citet{hinkle1988}, and in low-resolution interstellar absorption  in the far-IR ($\\nu_2$ bending mode) toward the Sgr B2 star forming region and IRC+10216 by \\citet{cernicharo2000}.  \\citet{giesen2001} discussed new laboratory data on the vibrational spectrum of \\cthree\\ in its low-frequency bending mode and re-visited the first identification of the $\\nu_2$ $R(2)$ line in absorption toward Sgr B2 \\citep{vanorden1995}. The abundance and excitation of \\cthree\\ with a large range of rotational temperatures in translucent clouds have also been determined convincingly \\citep{maier2001,roueff2002,oka2003,adamkovics2003} through observations at optical wavelengths. \\citet{galazutdinov2002} also demonstrated that the \\cthree\\ abundance is related neither to interstellar reddening nor to the intensities of diffuse interstellar bands.   DR21(OH) lies 2\\arcmin\\ North of the \\HII\\ region DR21, in the Cygnus X \\HII\\ complex. The distance to DR21(OH) has recently been accurately determined by trigonometric parallax measurements of its associated methanol masers as $1.50^{+0.08}_{-0.07}$~kpc \\citep{rygl2011}. Interferometric and high resolution single-dish observations both in continuum and molecular lines have shown multiple peaks in DR21(OH) \\citep{wilson1990,mangum1991,mangum1992,chandler1993}. \\citet{mangum1992} resolved the main DR21(OH) peak into two sources, MM1 and MM2, and MM2 into 2 sub-sources, MM2-A and MM2-B.  Synthesis imaging of NH$_3$ emission from the region has clearly resolved MM1 and MM2, having radial velocities $-4.1\\pm$0.3~\\kms\\ and $-0.7\\pm$0.3~\\kms, respectively \\citep{mangum1992}.  Both the star-forming cores MM1 \\& MM2 in DR21(OH) are young, having no visible \\HII\\ region and only very weak continuum emission at the centimeter wavelengths \\citep{argon2000}. Dust continuum observations suggest that MM1 is the brighter source ($L = 1.7\\times 10^4$~\\lsun; ZAMS B0V star) showing evidence of  star formation, whereas MM2 ($L = 1.2\\times 10^3$~\\lsun\\ ; early B star) though more massive is fainter and most likely at an earlier stage of evolution.  The region is characterized as a high-mass star-forming region due to the detection of centimeter and millimeter maser emission from numerous transitions e.g., H$_2$O, OH, CH$_3$OH \\citep{araya2009, fish2005, mangum1992, plambeck1990}.  MM1 also contains ground-state OH masers at 1.6~GHz and 6.7~GHz class II CH$_3$OH masers near the peak of the dust and centimeter continuum emission \\citep{argon2000,fish2005,rygl2011}.  In a recent interferometric study of the continuum emission at 1.4~mm from DR21(OH), \\citet{zapata2012} have resolved MM1 and MM2 into 9 compact sources. Five of the compact sources are associated with MM1 and four with MM2. Two of the compact sources associated with MM1 (SMA6 and SMA7) seem to show hot core activity.  \\citet{white2010} have observed multiple transitions of CO from 4--3 to 13--12 using SPIRE/Herschel to conclude the presence of two components at 80~K and 180~K from a rotation temperature analysis. From a LVG analysis these authors obtain T$_{\\rm kin}$ = 125~K and $n_{\\rm H_2}$ = 7$\\times 10^4$~\\cmcub.  We summarize selected properties of all the emission components in DR21(OH) as available from literature in Table ~\\ref{tab_srcprop}.  All values in Table~\\ref{tab_srcprop} correspond to a distance of 1.5~kpc.    The Heterodyne Instrument for the Far-Infrared \\citep[HIFI;][]{deGraauw2010} on board the Herschel Space Observatory \\citep{pilbratt2010}, with its broad frequency coverage, high sensitivity and spectral resolution has provided for the first time the opportunity for a systematic study of carbon chain molecules such as \\cthree\\ through probing several ro-vibrational lines at full spectral resolution.  The results presented here are a part of the PRISMAS (``PRobing InterStellar Molecules with Absorption line Studies\u201d) Key Program \\citep{gerin2010}.  The primary aim of the paper is to understand the \\cthree\\ formation mechanism. Observations of diffuse interstellar gas at optical and IR wavelengths have shown a strong correlation between the column densities of \\cthree\\ and C$_2$. This suggests that \\cthree\\ and C$_2$ are in the same chain of chemical reactions \\citep{oka2003}.  The formation routes for \\cthree\\ starting from C$_2$ also involved the production of CCH and \\cthreehtwo. However no transitions of C$_2$ are available at longer wavelengths (far-infrared and longer) which can be used to detect C$_2$ in dense star forming regions where \\cthree\\ is detected in absorption \\citep{mookerjea2010}.  Thus, to understand the formation pathway of \\cthree\\ we compare the  abundances of \\cthree, CCH, and \\cthreehtwo, with results of chemical models of dense star forming cores. ",
        "conclusions": "\\subsection{Lower gas phase abundance of \\cthree\\ in MM1}  The angular resolution of the spectroscopic observations used in this paper lies between 12 and 40\\arcsec, so that MM1 and MM2 are not resolved spatially. However, due to their markedly different velocities, MM1, MM2 and the envelope are easily discernible in the spectra. Interestingly, all \\cthree\\ spectra show primarily two velocity components associated with MM2 and the envelope, with absorption due to MM1 being detected only in $Q(4)$. In contrast, the chemically related species CCH and \\cthreehtwo\\  that we have considered here show components corresponding to only MM1 and the envelope.  Thus there is an intrinsic difference in the distribution of \\cthree\\ compared to the distribution of the other species. The abundance of \\cthree\\ in MM2 (1.2$\\times 10^{-9}$) is only twice that in the  envelope (6.3$\\times 10^{-10}$).  Based on Table~\\ref{tab_srcprop} the total column density of MM1 is $\\sim 60\\%$ of the column density of MM2. This suggests that if the relative abundances of \\cthree\\ in MM1 and MM2  were similar, the total column density of \\cthree\\ in MM1  would have been 2.5 times the column density in the envelope so that the absorption dips due to MM1 would be similar to or stronger than the dip due to the envelope.  It may be argued that the continuum backgrounds for the two components MM1 and the envelope may be different. However the continuum from MM1 is expected to be stronger than the continuum from the envelope.  All of the above arguments indicate that \\cthree\\ is preferentially depleted or destroyed in MM1.    One major difference between MM1 and MM2 lies in the fact that MM1 has already shown evidence for at least one high temperature hot core, whereas MM2, though more massive shows no such components.  The chemical network includes among others the reaction \\cthree\\ + H$_2$ $\\rightarrow$ C$_3$H + H \\citep{PFD88}, which becomes a primary destruction pathway for \\cthree\\ when the temperature exceeds $\\approx$ 80 K.  As a result, models predict that \\cthree\\ will be depleted in favor of \\cthreehtwo\\ in hot core conditions, as C$_3$H is also destroyed by reaction with H$_2$. Additionally, the chemical models presented here also show that the \\cthree\\ abundance is reduced at larger densities.  MM1 is definitely at a higher density than the envelope (Table~\\ref{tab_srcprop}), so that the reduced abundance of \\cthree\\ in MM1 relative to MM2 could also indicate that MM2 has a density lower than MM1. We note that this contradicts the densities derived from dust column densities (Table~\\ref{tab_srcprop}), which is not completely unexpected since the densities derived from the column densities are beam averaged and need not always reflect the local densities in a region.   However, an alternate possibility arises from the fact that the source intrinsic continuum opacity can re-fill the absorption to a significant extent, so that the column densities derived here are all primarily lower limits \\citep{mookerjea2010}.  It can thus be envisaged that in the case of MM1, which has a stronger continuum emission, it is a combination of the geometry of the source and the re-filling of the absorption dip which reduces the ``visibility\" of \\cthree. A proper evaluation of the effect  of the continuum opacity on the absorption depth is possible only by constructing a complete radiative transfer model of the entire region with accurate temperature and density profiles, which are not yet available.  Finally, since MM2 is estimated to be younger than MM1 and \\cthreehtwo, and CCH emission from MM2 are not detected, the possibility that  the observed \\cthree\\ abundance in MM2 is due to the pre-warm-up gas-phase chemistry can also not be ruled out.  \\subsection{Comparison with \\cthree\\ in diffuse clouds}  \\cthree\\ has been found to have an almost one-to-one correlation with C$_2$ in diffuse clouds \\citep{roueff2002}. Further, the abundance of \\cthree\\ in diffuse clouds, (3--6)$\\times 10^{-9}$ \\citep{roueff2002}, is ten times larger than the abundance of \\cthree\\ derived in the dense star forming environment of DR21(OH).  \\citet{oka2003} had explained the observed correlation between C$_2$ and C$_3$ in diffuse clouds, in terms of a direct pathway of formation of \\cthree\\ from C$_2$ (see their Fig. 4).  In the previous section, we examined the formation of C$_3$ during a warm-up period, and noted a relatively larger abundance of this species prior to $t=10^5$~yr, the so-called ``early time\".  We now explore the chemistry of \\cthree\\ in diffuse regions and during the early time period in dense gas regions, to examine the differences (if any).  The first step of the transformation of C$_2$ to C$_3$ in diffuse clouds is the photoionization of C$_2$ to C$_2^+$ \\citep{oka2003}. In contrast, the models for the dense star forming regions assume darker ($A_V = 10$) conditions, so photoionization is not an effective process.  In dense models, the process begins with the radiative association of C$^+$ + H$_2$ $\\rightarrow$ CH$_2^+$. Next, CH$_2^+$ reacts with H$_2$ to form CH$_3^+$.  This is followed by the dissociative recombination of CH$_3^+$ and CH$_2^+$ as competitive reactions to form CH, and then an ion-neutral  reaction with C$^+$ to form C$_2^+$.  As shown in Fig.~4 of \\citep{oka2003}, the pathway then follows along the upper (C$_2$H) branch to form C$_3$.  The C$_3$ is destroyed by processes that form C$_4$, C$_4^+$, and C$_5$, which react further to primarily form C$_3$ in a cyclic process.  The abundance of C$_3$ grows until C$^+$ is depleted at $t\\sim 10$--100~yr, depending on the density, after which it reaches a plateau until the C$_3$ is accreted onto grain surfaces. Dissociative recombination reactions with C$_2$H$^+$, C$_2$H$_2^+$, and C$_3^+$ along the pathway form C$_2$. The C$_2$ then reacts with O to form CO, and as a result, the abundance falls to $X(\\rm C_2) \\sim 10^{-9}$, far below the $X(\\rm C_2): X(\\rm C_3)=40:1$ measured by \\citet{oka2003} for diffuse clouds.   Thus, C$_2$ is not the starting point for C$_3$ formation in dense cloud models at early times, and also does not undergo a cycling process like that of C$_3$.  As a result, the correlation between C$_2$ and C$_3$ abundance is not predicted by the dark cloud models. Rather the C$_2$ abundance evolution more closely follows the evolution of CCH from early times until the warm-up begins than that of C$_3$.  This comparison indicates that the chemistry of \\cthree\\ formation in the envelope of DR21(OH) is definitely not the same as in diffuse clouds. Further, the early time abundance of C$_3$ for dense regions exceeds the observed abundance by three orders of magnitude, effectively excluding this formation route as well.   \\subsection{Warm Carbon Chain Chemistry in the envelope}  Comparison of the present model results and observations indicates that the observed composition can be simulated with ion-molecule chemistry following a moderate warm-up to 30-50~K. This also means that \\cthree\\ formation in star forming regions can be explained in terms of the moderate temperature (30~K) gas-phase chemical reactions starting from CH$_4$ evaporated from the grain mantle only, and does not require photochemistry of PAH molecules.  The observed abundances of C$_3$ are consistent with formation during a warm-up model, or ``Warm Carbon Chain Chemistry'' (WCCC) \\citep{sea07}, and related to the abundances of CCH, \\cthreehtwo. This suggests that the envelope conditions are more like those of ``lukewarm corinos'' surrounding low mass protostars than those of hot cores. In other words the \\cthree\\ is not residing in a hot core.  The abundances of the additional species chemically related to \\cthree, which we consider for the case of DR21(OH) provided further insight regarding the envelope conditions.  Table~\\ref{tab_obsabund} summarizes the observed abundances of the different species in DR21(OH) as well as those available in the literature for two types of sources: the lukewarm corinos  L1527 and B228 and the hot corino IRAS 16293-2422.  We find that while \\cthree\\ has not been observed in any of the other sources, the abundances of the other three species in DR21(OH) are closer to the abundances observed in lukewarm corinos. Further, the chemical model that best represents the observed abundances corresponds to a $T_{\\rm max}$ of 30--50~K, a temperature that is much lower than the values expected in hot cores/corinos. This suggests that the envelope of DR21(OH) is chemically closer to a WCCC region such as a lukewarm corino envelope than to a hot core. Further, the pre-depletion \\cthree\\ abundance (at $t\\leq 10^4$~yr) exceeds the observed values by a large factor. This implies that the \\cthree\\ abundance is primarily maintained via ion-molecule chemistry in the gas phase after CH$_4$ is desorbed from the surface of the dust grains, and does not have a one-to-one correlation with C$_2$ as found in diffuse clouds.  \\subsection{Chemical \\& Dynamical Ages of the Region}  The chemistry of the envelope surrounding the embedded sources MM1 and MM2 in DR21(OH) appears to have occurred over a period between (0.7--3)~Myr  including a cold period, warm-up, and extended time at $T = 30$ or 50 K.  This time period appears to be on the longer side of the evolutionary timescales of massive star forming cores. As detailed by \\citet{csengeri2011} the life time of the gas in  massive dense cores is determined by the crossing times or the local free-fall times for the molecular dense cores which is (5--7)$\\times 10^4$~yr.  However the large scale flows associated with Cygnus X (and DR21 in particular) are massive enough to continuously replenish the mass of these cores allowing them to remain ``active\" for a much longer period of time and the dynamical time of the most massive sub-filament in DR21 is $\\sim 2$~Myr \\citep{schneider2010}. Thus the chemical age of $\\sim 1$~Myr is consistent with the dynamical age of the region."
    },
    "1208/1208.4602_arXiv.txt": {
        "abstract": "{ Recent kinematical constraints on the internal densities of the Milky Way's dwarf satellites have revealed a discrepancy with the subhalo populations of simulated Galaxy-scale halos in the standard cold dark matter model of hierarchical structure formation. In particular, the Via Lactea II and Aquarius simulations both have large subhalos with internal densities that are larger than the constraints inferred for any Milky Way dwarf satellites. This has been dubbed the ``too big to fail'' problem, with reference to the improbability of large and invisible companions existing in the Galactic environment. In this paper, we argue that both the Milky Way observations and simulated subhalos are consistent with the predictions of the standard model for structure formation.  Specifically, we show that there is significant variation in the properties of subhalos among distinct host halos of fixed mass and suggest that this can reasonably account for the deficit of dense satellites in the Milky Way. We exploit well-tested analytic techniques to predict the properties in a large sample of distinct host halos with a variety of masses spanning the range expected of the Galactic halo.  Such techniques render the problem of estimating the variance in subhalo properties computationally feasible.  The analytic model produces subhalo populations consistent with both Via Lactea II and Aquarius, and our results suggest that natural variation in subhalo properties suffices to explain the discrepancy between Milky Way satellite kinematics and these numerical simulations.  At least $\\sim10\\%$ of Milky Way-sized halos host subhalo populations for which there is no ``too big to fail'' problem, even when the host halo mass is as large as $M_{\\mathrm{host}} = 10^{12.2}\\,h^{-1}$~M$_{\\odot}$.  Follow-up studies consisting of high-resolution simulations of a large number of Milky Way-sized hosts are necessary to confirm our predictions.  In the absence of such efforts, the ``too big to fail'' problem does not appear to be a significant challenge to the standard model of cold dark matter halos undergoing hierarchical formation. } ",
        "introduction": "In the standard paradigm describing structure growth in our universe, galaxies such as the Milky Way form within halos of dark matter \\cite{white_rees78,blumenthal_etal84}. This formation is hierarchical in the sense that small, self-bound clumps of dark matter collapse first into roughly virialized objects, subsequently merging at the nodes of sheets and filaments of dark matter to form ever larger halos.  Most of the merging halos are disrupted within a few dynamical times and are no longer recognizable as distinct objects within the larger halo; however, some survive to the present day within the virialized regions of their host halos. The properties of these dark subhalos and their luminous components, such as the satellite galaxies of the Milky Way, are probes of the faintest substructure surrounding the Milky Way and may serve as sensitive testbeds of cosmological structure formation \\cite{liddle_kamionkowski00,zb03,tollerud_etal08,bullock_etal10,font_etal11,rashkov_etal12}.  High-resolution simulations of $\\Lambda$CDM cosmological evolution have been able to model the properties of subhalos within host halos similar to the Milky Way.  However, the computational expense involved in such an undertaking has necessarily limited the number of experiments in which dwarf-sized satellites can be easily resolved at the ultra-faint scale ($L_{\\mathrm{dwarf}} \\gtrsim 10^5$~L$_{\\odot}$) to less than a dozen (taking as representative the set of Aquarius, GHALO, and Via Lactea II simulations in refs.~\\cite{springel_etal08}~\\cite{stadel_etal09}, and \\cite{diemand_etal07}, respectively).  Consequently, the diversity of subhalo characteristics and their natural variation among different host halos at fixed mass, as well as the details of those properties' dependence on halo mass, have only been subject to a relatively small number of studies.  Recent theoretical results in the context of Milky Way dwarf galaxy constraints suggest that the densest, most massive subhalos found in the Via Lactea II and Aquarius simulations have properties that are not commensurable with those of observed satellites in the Local Group \\cite{BK_etal11}. This builds upon the early indications and possible cosmological implications of such a mismatch explored in ref.~\\cite{zb03}, and raises the question of whether the densest, most massive dwarf galaxies predicted by simulations are present in the real Milky Way environment. If these few largest halos are present, then why are they not luminous like their slightly less massive counterparts?  In a properly cosmological context, it would seem that these putative satellites are ``too big to fail\", meaning that they are sufficiently more massive than the other Milky Way satellites that it is difficult to construct models in which these satellites remain invisible \\cite{BK_etal11}. In such a scenario, these subhalos are either dense and somehow dark relics \\cite{BK_etal11}, or have had their baryonic content scoured via early feedback effects \\cite{diCintio_etal11,BK_etal12,zolotov_etal12}.  An alternative to these astrophysical solutions would conclude that these few most massive subhalos are not to be found and that this discrepancy is truly reflective of a critical failure in the $\\Lambda$CDM structure formation paradigm (as proposed, for example, in the self-interacting CDM model of \\cite{vogelsberger_etal12}).  In this work, we explore another possible solution to the ``too big to fail\" problem.  We use analytic models to argue that the natural, statistical variation in subhalo densities from one host halo to another is significant and can affect the theoretical interpretation of satellite galaxy data.  In particular, we suggest that the properties of the few densest, most massive Milky Way satellites do not pose a challenge to galaxy or structure formation theories.  Rather, we argue that the observed satellite properties are consistent with being randomly drawn from among the statistical variety of subhalo populations that can be realized in Milky Way-sized halos, with a probability that is not negligible.  Halo-to-halo variation in subhalo properties is a natural consequence of the hierarchical structure-formation model as subhalos merge to form the contemporary host halo through an infinitude of distinct mass accretion histories. Any particular Galactic-scale halo is attended by a subhalo population with properties that are functions of this assembly history and likely a number of other variables.  At present, we have some ability to constrain viable formation routes of the Milky Way.  For example, our constraints on several properties of the thin and dynamically-cold Galactic disk imply that the host halo is unlikely to have suffered a minor merger (with mass ratio $\\lesssim 10:1$) passing through the disk within the past 5-8 Gyr \\cite{hammer_etal07,purcell_etal09}, before the ongoing Sagittarius impact and the imminent absorption of the Magellanic Clouds end that quiescent history at the present day \\cite{law_etal10,purcell_etal11,besla_etal07,busha_etal11}. However, existing constraints are broad and cannot be used to limit effectively the properties of the subhalos that should surround the Milky Way.  Moreover, we do not have a comprehensive understanding of which variables are most important in determining the properties of the Milky Way subhalo population.  At the same time, we have only a poor handle on the Milky Way's total mass because stellar rotation curves only probe the central regions of the Milky Way halo and other probes yield a wide range of values \\cite{xue_etal08,li_white08,przybilla_etal10,busha_etal11}. Interestingly and unfortunately, this mass range is fraught with cosmologically-significant transitions that bear directly on galaxy formation \\cite{codis_etal12,pichon_etal11,guo_etal10}, and numerical investigations have suggested that host halo mass is the largest determinant of subhalo properties in this particular context, with gas cooling and other baryonic effects taking a secondary role \\cite{geen_etal12}.  In addition, a recent statistical exploration of subhalo abundances in a simulated $\\Lambda$CDM universe has illustrated the sensitivity of satellite populations to host halo mass, finding that Galaxy-sized halos with mass $M_{\\mathrm{host}} \\lesssim 10^{12} $~M$_{\\odot}$ are much more likely to have only three large companions (analogous to the Large and Small Magellanic Clouds and the disrupting Sagittarius dwarf) at the present-day \\cite{wang_etal12}.  We complement this and other numerical efforts with our ability to probe host-mass parameter space, and most importantly the intrinsic variation among subhalo properties at fixed host-mass, in much greater detail due to the computational advantages of our analytic approach. In the remainder of the paper, we itemize our methods and results.  We describe the details of our formalism in \\S~\\ref{sec:methods}, including a restatement in \\S~\\ref{sec:toobig} of the ``too big to fail'' discrepancy in terms of a new variable describing central subhalo density.  We give our basic results, culminating in probability distributions for this density-proxy variable, in \\S~\\ref{sec:results}, reserving \\S~\\ref{sec:discuss} for discussion and interpretation in the context of future numerical experiments as well as achievable (and useful) observational constraints. ",
        "conclusions": "\\label{sec:discuss}  Using stellar kinematical data, ref.~\\cite{BK_etal11} compared constraints on the structural properties of the halos that host the dwarf satellite galaxies of the Milky Way with the structural parameters predicted by a number of numerical simulations of the formation of Milky Way-sized dark matter halos.  In doing so, they identified a discrepancy: too many numerical subhalos have $v_{\\mathrm{max}}$ values larger than would be expected for systems of their size, when set against these Galactic observations.  If $\\Lambda$CDM cosmological predictions and/or astrophysical models are not too badly wrong, then do massive invisible dwarfs exist around the Milky Way? How could such putative objects have failed to light up with star formation?  The authors of ref.~\\cite{BK_etal11} dubbed this issue the ``too big to fail\" problem.  Soon afterward, ref.~\\cite{wang_etal12} showed that one way to mitigate the ``too big to fail'' discrepancy is for the Milky Way to reside within a host dark matter halo at the lower-mass range of contemporary constraints on the size of the Milky Way halo.  Specifically, this investigation of the Millennium Simulation statistics found that for a threshold $v_{\\mathrm{max}} = 30$~km/s, fully $\\sim 40\\%$ of halos with host mass $M_{\\mathrm{host}} = 10^{12} M_{\\odot} =10^{11.85}~h^{-1} $~M$_{\\odot}$ host three or fewer satellites larger than the threshold (compared to the eight such satellites typical of the Aquarius halos).  For halos more massive than $M_{\\mathrm{host}} \\gtrsim 2 \\times 10^{12} M_{\\odot} =10^{12.15}~h^{-1} $~M$_{\\odot}$, this probability drops below $5\\%$ and quickly vanishes thereafter with increasing halo mass.  Our results are certainly in broad agreement with this conclusion, although the two predictions do not probe identical probability distributions; the summary statistic $\\Gamma$ involves $r_{\\mathrm{max}}$ as well as the $v_{\\mathrm{max}}$ distribution probed by ref.~\\cite{wang_etal12}. We therefore focus on the parameter $\\Gamma$ because it can directly address the degeneracies in halo properties permitted by the data, while noting that a simple translation of our $\\Gamma_{\\mathrm{max}} < 1$ criterion into $v_{\\mathrm{max}} < 30$~km/s results in $\\sim37\\%$ of $M_{\\mathrm{host}} = 10^{11.8} $~M$_{\\odot}$ hosts satisfying the constraints, in significant agreement with the sample analyzed by ref.~\\cite{wang_etal12}.  Proposed solutions to the  ``too big to fail\" problem, in looking for a means to reduce subhalo internal density, have ranged from the cosmological (if satellite centers are scoured by self-interacting dark matter as in ref.~\\cite{vogelsberger_etal12}) to the astrophysical (if dwarf galaxies sweep out dark mass via feedback effects as in refs.~\\cite{diCintio_etal11,BK_etal12,zolotov_etal12}). In this paper, we explore an additional way in which the ``too big to fail'' discrepancy may be mitigated.  In particular, we propose that the large variation in subhalo populations among different host halos can explain the dearth of large, dense subhalos orbiting the Milky Way without any making any adjustments to the host halo mass or accounting for baryonic feedback processes.  For a host halos at the high-mass end of the generally accepted range for the host halo of the Milky Way, $M \\approx 10^{12.2}\\, h^{-1}$~M$_{\\odot}$, we find that at least $\\sim 10\\%$ of all host halos would harbor a subhalo population consistent with the observed stellar kinematics of the Milky Way dwarfs.  This probability can be considerably higher if the three objects neglected in the \\cite{BK_etal11} analysis (Sagittarius and the Magellanic Clouds) exhibit any of the highest densities among the Milky Way subhalo population.  In formulating our comparison of theoretical predictions with observational constraints, we have introduced a new density-proxy parameter $\\Gamma$, that roughly runs perpendicular to the degeneracy between the subhalo structural parameters $v_{\\mathrm{max}}$ and $r_{\\mathrm{max}}$ exhibited by the data.  Our detailed results are statements about the approximate probability distribution of the parameter $\\Gamma$ among subhalos in different hosts.  Our study has not identified additional diagnostics that can be used to characterize the subset of host halo merger histories that produce satellite populations not subject to the ``too big to fail'' problem.  From Fig.~\\ref{fig:gamma}, we see that enforcing a recent epoch of relative quiescence on the host halo does not significantly increase probability for harboring an acceptable subhalo population.  The one exception to this conclusion is for the lowest host halo mass, $M_{\\mathrm{host}} = 10^{11.8}\\,h^{-1}$~M$_{\\odot}$.  The probability increase reflects the fact that less massive halos are more likely to have experienced minor mergers of a given mass ratio by the present day \\cite{stewart_etal08}, and this relative homogeneity among host halos also manifests in the cumulative velocity functions and average concentration parameters $c$ shown in Fig.~\\ref{fig:nvfn}, which do not vary outside the Poisson error on the mean for the $M_{\\mathrm{host}} = 10^{11.8}\\,h^{-1}$~M$_{\\odot}$ model.  Using the large variance of available merger histories to test the observational consistency of any particular realization's subhalo population could be an avenue of potential interest in an era of near-field cosmology that seeks to connect specific substructure in the Milky Way to characteristics of Galaxy-formation theory. Our rough test along those lines in the present work differentiates mainly between systems that {\\em have} undergone their allotted minor mergers versus those that {\\em have not yet} been impacted by high-mass subhalos; in a sense, we already know that the Galaxy is part of the latter class, due to the unusually-thin and dynamically-cold state of the stellar disk, and that in fact it is presently undergoing multiple minor mergers. However, to get a much more detailed handle on the assembly history of the Milky Way by inspecting analytic distributions of $\\Gamma$ and $\\Gamma_{\\mathrm{max}}$, one would require a reasonable estimate of $v_{\\mathrm{max}}$ for each individual satellite galaxy's host halo, and it is unclear what (if any) fundamental science could be done with such an effort, beyond refining the correspondence between numerical results, theoretical predictions, and observational constraints (which we argue here to be statistically sound, in contrast to the claims implicit to the ``too big to fail\" conjecture).  Fig.~\\ref{fig:nvfn} shows subtle distinctions between the properties of host systems at the two larger masses ($M_{\\mathrm{host}}=10^{12.0}\\,h^{-1}$~M$_{\\odot}$ and $M_{\\mathrm{host}}=10^{12.2}\\,h^{-1}$~M$_{\\odot}$) that pass the ``too big to fail'' test and those that do not.  Hosts halos of these masses that contain subhalo populations that are grossly consistent with the Milky Way satellite populations also generally have concentrations that are slightly higher than average, though this systematic offset is smaller than the dispersion in concentration at fixed mass \\cite{bullock_etal01,neto_etal07,maccio_etal07}. This weak trend is broadly consistent with the sense that the Milky Way may be an unusually small disk galaxy among similarly-sized systems, in terms of stellar mass as well as angular momentum \\cite{hammer_etal07}, as $\\Lambda$CDM models predict that disk size increases with decreasing host halo concentration \\cite{bullock_etal01}. This implication also appears to be consistent with recent measurements that correlate satellite concentration with host stellar mass and total halo mass, in nearby disk-dominated galaxies \\cite{cibinel_etal12}.  A possibly more dramatic feature of subhalo populations within host halos in our higher-mass samples that are not in conflict with Milky Way data is that they exhibit relatively steeper cumulative velocity functions, as shown in Fig.~\\ref{fig:nvfn}.  At best, this only slightly ameliorates the discrepancy that has come to be known as the ``missing satellite'' problem, if we take low $\\Gamma_{\\mathrm{max}}$ samples to represent possible Milky Way realizations.  However, as large-scale surveys of the Galactic environment complete the local census of dwarf satellites, the faint-end slope may probe this prediction of models that solve the ``too big to fail'' problem by exploiting the large variance in subhalo populations.  We note, however, that Fig.~\\ref{fig:nvfn} appears to indicate that halo mass and velocity-function slope may be somewhat degenerate, and therefore a full exploration of this parameter space by observations of substructure in extragalactic systems could be useful in discriminating between the competing effects we discern here.  Interestingly, the steeply-sloped realizations with low $\\Gamma_{\\mathrm{max}}$ also display a plateau beneath a value of unity in the subhalo velocity function.  This reflects the fact that some relatively rare Galaxy-sized halos may have a very massive companion with $v_{\\mathrm{max}} > 60$~km/s, with such a subhalo having necessarily fallen in recently (since dynamical friction destroys massive objects quickly).  We note that the Milky Way dwarf population mimics this behavior somewhat, within observational Poisson error of the analytic plateau, due to the presence of three large satellites with $v_{\\mathrm{max}} > 60$~km/s as well as the absence of any moderately-massive subhalos between $30 < v_{\\mathrm{max}} < 60$~km/s, supporting the conclusion that the low $\\Gamma_{\\mathrm{max}}$ description of the Galactic dwarfs is consistent with a flat velocity function in this intermediate range, and favors the presence of one or possibly more large systems with relatively higher $v_{\\mathrm{max}}$.  In the particular case of the Milky Way, this plateau occurs at three large satellites, due to the conjoined Magellanic Clouds and the currently-disrupting Sagittarius dwarf galaxy.  In summation, we have used analytic models of subhalo populations to argue that the absence of luminous Milky Way satellite galaxies residing in halos of high density may not indicate a problem with contemporary theories of structure formation or galaxy formation.  Rather, we have argued that the non-negligible variation in subhalo populations may reasonably account for this deficit.  We estimate that at least 10\\% of halos could be consistent with observations without any requirement that the host halo of the Milky Way have a mass toward the lower range of contemporary constraints.  There are several implications of this result.  First, it suggests that there is no need to consider relatively low-mass hosts for the Milky Way (though such may certainly be the case, and could be a large part of the solution).  Second, it shows that the ``too big to fail'' problem cannot be a statement of high statistical significance.  Further study is necessary to confirm or refute our argument.  In particular, we have, by necessity, used approximate techniques to predict the properties of dark matter subhalos so a large numerical simulation campaign will be necessary to test our predictions. In performing such a follow-up study, it must be borne in mind that selecting particular halos for high-resolution resimulation using accretion history or halo structural information \\cite{springel_etal08,diemand_etal07} can introduce biases; however, we have not been able to identify clear biases in our analytic models.  Furthermore, we have neglected the effects of baryons and these may need to be considered more carefully.  Finally, once a large set of simulated subhalo populations is available, it will be imperative to conduct a full statistical comparison of theoretical predictions with stellar kinematical data.  Such follow-up studies may, indeed refute our argument perhaps by identifying undiagnosed systematic errors in our methods. However, absent such detailed follow-up, the ``too big to fail'' problem is unlikely to pose a serious challenge to the standard, hierarchical cold dark matter model of structure growth; rather than being a potential wound to the predictive power of $\\Lambda$CDM, it may instead represent an exhibition of its productive variety."
    },
    "1208/1208.1966.txt": {
        "abstract": "\\begin{center} {\\bf \\large Abstract} \\\\ \\end{center}  The Arcminute Microkelvin Imager~(AMI)~\\citep{jon08} is a telescope specifically designed for high sensitivity measurements of low-surface-brightness features at cm-wavelength and has unique, important capabilities. It consists of two interferometer arrays operating over 13.5--18~GHz that image structures on scales of 0.5--10 arcmin with very low systematics.  The Small Array (AMI-SA; ten 3.7-m antennas) couples very well to Sunyaev-Zel'dovich~(SZ) features from galaxy clusters and to many Galactic features.  The Large Array~(AMI-LA; eight 13-m antennas) has a collecting area ten times that of the AMI-SA and longer baselines, crucially allowing the removal of the effects of confusing radio point sources from regions of low surface-brightness, extended emission. Moreover AMI provides fast, deep object surveying and allows monitoring of large numbers of objects. In this White Paper we review the new science --- both Galactic and extragalactic --- already achieved with AMI and outline the prospects for much more. ",
        "introduction": "\\label{sz}  \\subsection{Introduction}  Clusters of galaxies are the most massive bound structures in the cosmic web of large-scale structure.  They evolve through the linear regime from the density perturbations imaged in the primordial cosmic microwave background (CMB) and then continue to grow in the non-linear regime through in-fall and mergers.  As such, clusters have great potential for use in cosmology, both to investigate the aspects of the underlying cosmological model (e.g. $\\sigma_8$) and as cosmographic buoys to study the evolution of structure as a whole in the universe (see e.g. \\citet{2011ARA&A..49..409A}).  In order to connect clusters to the underlying cosmology, the critical parameter that must be measured is their mass. Various scaling relations have been proposed for mass determination, but the Comptonisation $Y$ parameter is now recognised as a highly-reliable mass indicator \\citep{2006ApJ...650..128K}. This is, of course, measured directly through SZ measurements, without the need for a model-dependent calculation from the data as is the case for X-rays.  The redshift-independence of SZ surface brightness offers the possibility of observing clusters right back to their epoch of formation; and the linear (as opposed to quadratic) dependence of the SZ signal on gas density allows imaging of the outer regions of the Intra Cluster Medium~(ICM) --- \\citet{2011ARA&A..49..409A} emphasise the importance of investigating these outer regions, which is challenging in X-rays, but straightforward in SZ. Combining SZ measurements with X-ray observations and optical-lensing data gives essential, new information.  In addition to their use for cosmology, clusters are of fundamental interest in their own right, harbouring a great deal of detailed ICM astrophysics (see e.g. \\citet{2011MNRAS.410.2446K}). For example, merger events are laboratories in which the nature of dark matter can be tested \\citep{2006ApJ...648L.109C}.  Imaging of detailed ICM structure is now becoming possible in both X-ray and SZ (e.g. \\citet{2010ApJ...716..739M,2011ApJ...734...10K}), allowing measurement of density and temperature over a range of scales.  A critical issue in understanding the ICM is the role of magnetic field. This could substantially stiffen the ICM equation of state, affect cluster dynamics and evolution, and lead to bias in many mass estimations. Little is known about ICM magnetic field, because the requisite measurements have been out of reach: however, the instrumental situation is now changing.  The initial science drivers for AMI were, via Sunyaev-Zel'dovich effect~(SZ) measurements, to constrain cosmology through conducting a blind survey for galaxy clusters; and to explore the physics of clusters and tie down their scaling relations and evolution through pointed observations towards known or candidate clusters. AMI has made significant advances in both of these areas.  \\subsection{Blind cluster SZ surveys}  We have surveyed 10 first-look fields to target sensitivities of $100\\,\\mu$Jy and $60\\,\\mu$Jy with AMI's Small (SA) and Large (LA) arrays respectively. Each candidate cluster identified through this analysis is then reobserved with deep, pointed observations to verify the detection.  In \\citet{tim10} we describe the first successful cluster detection from the first (blind) field to be completed (see Figure~\\ref{firstblind}).  This distant, remarkably extended system, probably a merger, is detected at very high significance by AMI in SZ. We calculate that the cluster mass limit to which our survey will be sensitive as $2\\times10^{14} M_\\odot$ by scaling from the reported detection.  A full Bayesian analysis \\citep{2012MNRAS.421.1136A} of the first 10 fields is underway to detect clusters, determine their significance and place constraints on cosmological parameters, e.g.  $\\sigma_8$. In addition, the AMI-LA observations of our fields have allowed us to measure the radio source counts at 15.7~GHz down to 0.5~mJy \\citep{2011MNRAS.415.2699A,2011MNRAS.415.2708A}; this 10C survey is the deepest radio survey of any significant extent above 1.4~GHz. New programme of observations will push the source counts a further factor of three deeper. Preliminary results suggest the detection of a change in the slope of the source counts and a change in the spectral index distribution, both of which indicate a change in the radio source population.  AMI SZ-surveying and analysis is now routine and has been demonstrated to detect clusters down to a limiting mass of $2\\times10^{14} M_\\odot$. Therefore there is a great deal of scientific potential in performing complementary surveys towards fields observed in other wavebands.  For example, the AMI blind surveys have now been extended to include additional fields toward the Lockmann Hole and ELAIS N1 regions which have been surveyed deeply by UKIDSS. This ongoing project will allow cross-calibration of IR and SZ scaling relationships and investigation of cluster selection functions of the two methods.  This project will also provide a legacy survey of these well-studied regions.  Another example of a future collaboration is with Battye et al.~on the SuperCLASS legacy programme (http://www.e-merlin.ac.uk/legacy/projects/superclass.html), which has recently been allocated 862 hours of eMERLIN observing time. The aim of this programme is to measure weak gravitation lensing through radio imaging towards a supercluster field. Combining these data with SZ measurements from AMI will set very powerful joint constraints on parameters of baryonic and dark matter structures in this field.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=11.0cm]{SA_pointed_a_sub_T323M_box} \\caption{Map of the SZ effect in the first cluster system to be discovered in the AMI blind survey. positive contours are solid and negative ones dashed; contour levels are in units of $\\sigma$. Crosses (vertical and diagonal) indicate the position of point radio sources that have been subtracted from the data. Squares show the best fit positions for the clusters.  The synthesised beam is indicated in the lower-left corner.} \\label{firstblind} \\end{center} \\end{figure}  \\subsection{Pointed SZ observations}  AMI has been used for a wide range of SZ projects targeting fields of interest revealed in other wavebands \\citep{ami06,tash11,carmen11a,carmen11b,jon11}.  We have demonstrated the capabilities of AMI through observations of a variety of different known clusters \\citep{jon11}. We have also used these observations to validate our fully Bayesian analysis tool, {\\sc McAdam}~\\citep{2009MNRAS.398.2049F}, which uses nested sampling to fit various cluster parameterisations efficiently to our data, in particular generating a robust estimate of the key observable of cluster mass out to the virial radius.  AMI has been used for a variety of different pointed SZ programmes: to study the astrophysics of individual clusters of particular interest e.g. the bullet-like cluster Abell 2146 (see Figure~\\ref{a2146}) and the Corona Borealis supercluster; to produce joint constraints on cluster models once combined with optical lensing data \\citep{tash12}; for statistical analyses of scaling relations through observations towards well-defined samples of clusters drawn from X-ray surveys (e.g. the LOCUSS, BCS and MACS samples).  We now detail one particular example as follows.  AMI observations of an X-ray luminosity-limited subset of 19 LoCuSS clusters (which have $0.142 \\le z \\le 0.295$ to minimise cosmic evolution) have yielded 16 SZ detections with significances (of the peaks alone) between $5 \\sigma$ and $23 \\sigma$ (see Figure~\\ref{locuss} for SZ maps overlaid on X-ray for four of these clusters). The SZ images span scales out to the cluster virial radii (which very few X-ray images do), and show a larger range of morphology than that seen in published ACT, SPT or CARMA/SZA images (which tend not to reach $r_{\\mathrm{virial}}$) or in Planck images (which do not have enough angular resolution).  Under reasonable assumptions, the cluster plasma temperature can be derived from AMI data alone \\citet{carmen11b} --- unlike the traditional method, no X-ray information is required. In the 10 cases where deep \\textit{Chandra} or Suzaku images are available, allowing X-ray temperatures to be measured out to around 0.5$r_{\\mathrm{virial}}$ (rather than the more common 0.1--0.2$r_{\\mathrm{virial}}$), it is clear that there is decent correspondence between these $T_{X, 0.5\\mathrm{virial}}$-values and the $T_{AMI, 0.5-1\\mathrm{virial}}$ values from SZ, but with some strong outliers that have high $T_{\\mathrm{X}}$ compared with $T_{\\mathrm{AMI}}$. These outliers are the strong mergers. AMI easily probes the outer parts of clusters where most of the baryons and dark matter exist; these parts are hard for X-ray instruments to detect even at moderate $z$. At moderate $z$, the combination of AMI and \\textit{Chandra}/Suzaku is thus very powerful for understanding how clusters work.  There are several ongoing collaborations that will allow joint analysis of our AMI images with data from other wavebands, as summarised below.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=11.0cm]{A2146_mcsub_sigcont_312200} \\caption{AMI map of the SZ effect in the Bullet-like cluster A2146. Crosses indicate the positions of subtracted radio sources. The ellipse lower-left indicated the synthesised beam.} \\label{a2146} \\end{center} \\end{figure}  \\begin{figure}[t] \\begin{center} \\includegraphics[width=7.5cm,height=6.5cm,clip=,angle=0.]{./A611_xray}\\qquad\\includegraphics[width=7.5cm,height=6.5cm,clip=,angle=0.]{./A990gimp_nobar_ndec}\\\\ \\includegraphics[width=7.5cm,height=6.5cm,clip=,angle=0.]{./A2111_xraygimp}\\qquad \\includegraphics[width=7.5cm,height=6.5cm,clip=,angle=0.]{./A2218_xraygimp} \\caption{Maps of four clusters from the LoCuSS sample; from top left: A611, A990, A2111, A2218. Contours show the SZ decrement measured by the AMI-SA after source subtraction, while the greyscale is the X-ray image smoothed {\\it Chandra} data except for A990 which is {\\sc \\textit{ROSAT} HRI}).\\label{locuss}} \\end{center} \\end{figure}   \\begin{itemize}  \\item {\\bf MACS sample} This is a high-redshift sample and so offers the possibility of investigating potential evolution in cluster properties. Deep \\textit{Chandra} data and X-ray temperatures can be used to constrain the priors for our {\\sc McAdam} analysis.  \\item{\\bf The Canadian Cluster Comparison Project~(CCCP) sample}. The CCCP team have CFHT lensing observations towards a large sample of X-ray selected clusters; many have been observed with \\textit{Chandra} and \\textit{XMM}.  \\item {\\bf The CLASH sample}. The CLASH team has lensing data from \\textit{HST} and Subaru of a sample of dynamically-relaxed, massive clusters and this will allow us to derive stringent constraints on dark matter and hot baryons in the representative clusters.  \\item {\\bf Serendipitous \\textit{XMM} cluster detections} \\textit{XMM} images and temperatures will be available for combination with the AMI data.  \\end{itemize}   There continue to be huge opportunities for new science from pointed AMI SZ observations as new cluster samples become available and for joint analyses with other data sets to explore the physics of the intracluster gas of particular systems. In addition, there are possibilities for entirely new AMI SZ science areas, such as attempts to detect the SZ effect from the plasma in filaments of the cosmic web, and combination of a AMI-derived gas model with rotation-measure observations from the new generation of radio observatories such as JVLA and MeerKAT to measure cluster magnetic fields. One of the most exciting opportunities in the short term is support of the \\textit{Planck} satellite, which is now providing the first all-sky SZ cluster survey. This is discussed in the section~\\ref{ami-planck}.  \\subsection{Comparison of AMI to other SZ telescopes}  The South Pole Telescope (SPT) and the Atacama Cosmology Telescope (ACT) are multipixel bolometric telescopes specifically designed for SZ work. They are primarily survey instruments and have already produced extremely exciting results (see e.g. \\citet{2012arXiv1203.5775R} for SPT and \\citet{2011ApJ...737...61M} for ACT). Both SPT and ACT are conducting very wide field, relatively shallow surveys and so are detecting the rare, most massive clusters in the Universe, while AMI's survey strategy is to spend longer integration times on a smaller area and so sample typical clusters.  AMI is therefore able to probe the cluster mass function significantly deeper and so the three telescopes produce useful complementary information on structure formation. AMI is an excellent instrument for pointed follow-up observations on samples of galaxy clusters from other instruments since its 20~arcmin field-of-view matches the typical size of $z>0.15$ clusters very well; SPT and ACT's very wide fields-of-view mean that their strength is in survey observations. Finally, both ACT and SPT are located in the southern hemisphere with the result that they have little or no (respectively) coverage of the northern sky visible to AMI.  The SZA is now part of the CARMA telescope and so is no longer a dedicated SZ instrument. It is similar in many ways to AMI, but AMI has two key differences which are important for SZ science: \\begin{itemize} \\item AMI has access to a factor of two times shorter baselines (in terms of wavelength, i.e. in the uv-plane), and so samples structure out to 10 arcmin, allowing mapping of the gas out to the virial radius in galaxy clusters between z=0.15-2. \\item The AMI-LA has ten times the collecting area (and so ten times the flux sensitivity) of the AMI-SA and a good coverage of the uv-plane. This gives very high sensitivity, high image fidelity mapping of the radio point sources, which are the key contaminant and dominant systematic for SZ work. For comparison, SZA configures two of its eight antennas as outriggers for source subtraction, giving comparatively poor flux sensitivity and a single narrow strip of coverage in the uv-plane, resulting in a poor point-spread function and so poor image fidelity. \\end{itemize} ",
        "conclusions": ""
    },
    "1208/1208.1950_arXiv.txt": {
        "abstract": "This thesis is the result of my work as research fellow at \\href{http://www.iasf-milano.inaf.it}{IASF-MI}, Milan section of the Istituto di Astrofisica Spaziale e Fisica Cosmica, part of INAF, Istituto Nazionale di Astrofisica. This work started in January 2006 in the context of the PhD school program in Astrophysics held at the \\href{http://www.fisica.unimi.it}{Physics Department of Universita' degli Studi di Milano} under the supervision of Aniello Mennella.  The main topic of my work is the software modelling of the Low Frequency Instrument (LFI) radiometers. The LFI is one of the two instruments on-board the European Space Agency Planck Mission for high precision measurements of the anisotropies of the Cosmic Microwave Background (CMB).  I was also selected to participate at the International Doctorate in Antiparticles Physics, \\href{http://www.fe.infn.it/idapp}{IDAPP}. IDAPP is funded by the Italian Ministry of University and Research (MIUR) and coordinated by Giovanni Fiorentini (Universita' di Ferrara) with the objective of supporting the growing collaboration between the Astrophysics and Particles Physics communities. It is an international program in collaboration with the Paris PhD school, involving Paris VI, VII and XI Universities, leading to a double French-Italian doctoral degree title.  My work was performed with the co-tutoring of Jean-Michel Lamarre, Instrument Scientist of the High Frequency Instrument (HFI), the bolometric instrument on-board Planck. Thanks to this collaboration I had the opportunity to work with the HFI team for four months at the Paris Observatory, so that the focus of my activity was broadened and included the study of cross-correlation between HFI and LFI data. Planck is the first CMB mission to have on-board the same satellite very different detection technologies, which is a key element for controlling systematic effects and improve measurements quality.  The thesis is organised in four chapters: \\begin{enumerate} \\item a short introduction focused on state-of-art CMB phenomenology \\item a chapter about the Planck mission mainly focused on the LFI instrument \\item a chapter about software modelling of the LFI radiometers which includes a detailed description of the LFI instrument. Here I also discuss the experimental data available from the measurements campaigns on radiometer components, the model implementation and its validation against the frequency response measurements \\item a chapter about the satellite thermal environment, with particular reference to  the stage cooled at 4K, which is of key importance for both instruments. In this chapter I show the result of the analysis of the propagation of temperature fluctuations through the HFI. \\item a chapter about cross-correlation of HFI and LFI data. In this chapter I describe the implemention of data analysis sessions in the KST data visualization software with the purpose of simplifying and standardsing the cross-correlation analysis \\end{enumerate} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.1999_arXiv.txt": {
        "abstract": "We establish that the light Higgs boson mass in the context of the No-Scale Flipped $SU(5)$ GUT with TeV scale vector-like matter multiplets (flippons) is consistent with $m_h = 125.5\\pm0.5$ GeV in the region of the best supersymmetry (SUSY) spectrum fit to low statistics data excesses observed by ATLAS in multijet and light stop 5 \\fb SUSY searches at the LHC7.  Simultaneous satisfaction of these disparate goals is achieved by employing a minor decrease in the $SU(5)$ partial unification scale $M_{32}$ to lower the flippon mass, inducing a larger Higgs boson mass shift from the flippon loops. The reduction in $M_{32}$, which is facilitated by a phenomenologically favorable reduction of the low-energy strong coupling constant, moreover suggests an imminently observable ${(e\\vert\\mu)}^{\\!+}\\! \\pi^0$ proton decay with a central value time scale of $1.7\\times10^{34}$~years.  At the same point in the model space, we find a lightest neutralino mass of $m_{\\chi} = 145$ GeV, which is suitable for the production of 130 GeV monochromatic gamma-rays through annihilations yielding associated $Z$-bosons; a signal with this energy signature has been identified within observations of the galactic center by the FERMI-LAT Space Telescope. In conjunction with direct correlations to the fate of the ATLAS multijet and light stop production channels presently being tested at the LHC8, we suggest that the reality of a 125.5~GeV Higgs boson affords a particularly rich company of specific and imminently testable associated observables. ",
        "introduction": "The recent combined discovery of the CP-even light Higgs boson around $m_h = 125$-$126$~GeV by the ATLAS~\\cite{ATLAS:2012gk}, CMS~\\cite{CMS:2012gu}, and CDF/D0~\\cite{Aaltonen:2012qt} Collaborations has sparked detailed examinations into what regions of the SUSY parameter space remain viable, and which are very constrained, if not outright excluded.  Motivated by these new precision experimental measurements, we endeavor here to commensurately enhance the precision of prior theoretical estimates~\\cite{Li:2011xg,Li:2011ab} for the range of Higgs boson masses consistent with the No-Scale Flipped $SU(5)$ model with vector-like matter (flippons)~\\cite{Li:2010ws,Li:2010mi,Li:2010uu, Li:2011dw,Maxin:2011hy,Li:2011hr,Li:2011xu,Li:2011in,Li:2011gh,Li:2011rp, Li:2011fu,Li:2011xg,Li:2011ex,Li:2011av,Li:2011ab,Li:2012hm,Li:2012tr,Li:2012ix,Li:2012yd,Li:2012qv}, dubbed \\fsu5 for short.  The mutual consistency of a $125$-$126$~GeV Higgs boson mass with an explanation for the tantalizing positive excesses~\\cite{Li:2012tr,Li:2012ix} observed at low statistics by ATLAS~\\cite{ATLAS-CONF-2012-033,ATLAS-CONF-2012-037} and CMS~\\cite{SUS-12-002} in the $\\sqrt{s} = 7$ TeV SUSY search dramatically narrows the \\fsu5 model space, with interesting and specific implications for ongoing collider, indirect dark matter detection, and proton decay experiments.  The central mass peak of the new particle observed at the LHC exists at 125.3 GeV and 126.0 GeV, according respectively to the CMS~\\cite{CMS:2012gu} and ATLAS~\\cite{ATLAS:2012gk} collaborations. In each case, statistical and systematic errors of about half a GeV are suggested. Lacking an official combination of the two experiments' statistics, a more broad treatment might follow something akin to the proposal $m_h = 125.0 \\pm 1.0~{\\rm (exp)} \\pm 1.5~{\\rm (theory)}$~GeV of Ref.~\\cite{Buchmueller:2011ab}, noting in particular that our capacity to computationally model the predicted Higgs value is likewise imprecise. However, it is our interest in the present work to ascertain the restrictions that a Higgs mass very close to the average of centrally reported values would place on an otherwise successful model.  Insomuch as this mass is somewhat heavier than what may be comfortably achieved in typical GUT models without resorting to very heavy sparticles, we are interested in demonstrating a counter-example satisfying that vital emerging constraint which does not resort to taking extremities in the admissible lower bound.  In particular, we presently investigate the possibility that a large portion of the burden could be absorbed (in conjunction with the described flippon loops) by a lowering of the strong coupling $\\alpha_{\\rm s}$ at $M_{\\rm Z}$, an accommodation to which the flipped $SU(5)$ GUT is particularly well historically adapted~\\cite{Ellis:1995at}.  The Higgs mass constraint $m_h = 125.5\\pm0.5$ adopted for this work is thus purposefully rather strict. ",
        "conclusions": ""
    },
    "1208/1208.2097_arXiv.txt": {
        "abstract": "We present a physically motivated semi-analytic model to understand the clustering of high redshift Lyman break galaxies (LBGs). We show that the model parameters constrained by the observed luminosity function, can be used to predict large scale (i.e $\\theta\\ge 80''$) bias and angular correlation function of galaxies. These predictions are shown to reproduce the observations remarkably well. We then adopt these model parameters to calculate the halo occupation distribution (HOD) using the conditional mass function. The halo model using this HOD is shown to provide a reasonably good fit to the observed clustering of LBGs at both large and small ($\\theta < 10''$) angular scales for the whole range of $z=3-5$ and limiting magnitudes. However, our models underpredict the clustering amplitude at intermediate angular scales, where quasi-linear effects are important. The average mass of halos contributing to the observed clustering is found to be $ 6.2 \\times 10^{11}$ M$_\\odot$ and the characteristic mass of a parent halo hosting  satellite galaxies is $1.2 \\times 10^{12}$ M$_\\odot$ for a limiting absolute magnitude of $-20.5$ at $z=4$. For a given threshold luminosity these masses decrease with increasing $z$ and at any given $z$ these are found to increase with increasing value of threshold luminosity. Our physical model for the HOD suggests that approximately $40\\%$ of the halos above a minimum mass $M_{min}$, can host detectable central galaxies and about $5-10\\%$ of these halos are likely to also host a detectable satellite. These satellites form typically a dynamical timescale prior to the formation of the parent halo. The small angular scale clustering is mainly due to central-satellite pairs rather than few large clusters. It is quite sensitive to changes in the duration of star formation in a halo and hence could provide a probe of this quantity. The present data favor star formation in a halo lasting typically for a few dynamical time-scales, with 50\\% of stars formed in a time $T \\sim 300-500$ Myr for dark matter halos that collapse in the redshift range of $5.5-3.5$. Our models also reproduce different known trends between parameters related to star formation. ",
        "introduction": "Over the past decade there has been a growing wealth of observations probing the properties of high redshift galaxies. Various surveys, using the Lyman break color selection technique \\citep{madau_96, steidel_96_1,steidel_98,steidel_98_1}, have detected a substantial number of high redshift galaxies, up to $z\\sim 10$. This has resulted in reasonably good estimates of luminosity functions (LF) of these Lyman break galaxies (LBG) up to $z\\sim 8$ \\citep{bouwens_07_LF_z46,bouwens_07_LF_z710, reddy_08_LF,vanderburg_10_LF} and also LBG clustering up to $z \\sim 5$ \\citep{giavalisco_dickinson_01,porciani_giavalisco_02,ouchi_04_acf, adelberger_steidel_05,ouchi_hamana_05_acf,kashikawa_06_acf, lee_giavalisco_06_acf,hildebrandt_09_acf,savoy_11_acf,bielby_11_acf}. It is important to explain these observations and understand their implications for galaxy formation.  In the hierarchical model of structure formation galaxies form in virialized dark matter halos. These inturn result from the growth and gravitational collapse of initial Gaussian density perturbations. Thus the statistical properties of galaxies are determined by the statistics of the parent halo population, given a model for how stars form inside these halos. The properties of dark matter halos are quite well understood using N-body simulations \\citep{springel_white_05} and analytical models like the halo model of large scale structure \\citep{cooray_sheth_02}. These approaches provide the abundance, spatial distribution and merger history of dark matter halos. Numerical simulations also suggest a possible universal dark matter halo density profile, NFW profile \\citep{NFW_97}. Given the above inputs on dark matter halo properties and a specific model of galaxy formation inside these halos, it is possible to explain the two major observables of galaxies, their luminosity function and clustering. In addition, such models can throw light on the complex physics of galaxy formation, such as rate and duration of star formation, feedback mechanisms etc.  There has been extensive modelling of the luminosity functions and clustering of galaxies at low redshifts \\citep{somerville_99,yang_mo_03, zheng_zehavi_09,zehavi_zheng_11}. Several studies on understanding the LF of high redshift LBGs have also been carried out \\citep{somerville_01, benson_bower_03,stark_loeb_07,khochfar_silk_07}. We have been exploring physically motivated semi-analytic models of galaxy formation to understand the LFs of LBGs and Lyman-$\\alpha$ emitters, galactic winds and reionization of the intergalactic medium \\citep{samui_07,samui_08,samui_09,samui_09_lae,samui_10}. In our model, the luminosity of any galaxy is obtained from a physical model of star formation rate which depends on the mass and age of the hosting halo. We then combine this information with the formation rate of dark matter halos to obtain the LF of LBGs at various redshifts (see section 2 for details). These models have reproduced the LFs of high $z$ LBGs from $z=3-7$ reasonably well. In addition they constrain the efficiency of star formation and its duration in LBGs, and have also been used to set tight limits on the neutrino mass \\citep{charles_11}. We now wish to examine if our simple physical model for galaxy formation, combined with the halo model, can also explain the clustering of the high $z$ LBGs.   Semi-analytical models of clustering involve giving a prescription for how many galaxies of different luminosities occupy a dark matter halo of a given mass. This is called the Halo occupation distribution (HOD) and is usually given in a parametrized form \\citep{jing_98,seljak_00_HOD,scoccimarro_sheth_01,bullock_wechsler_02_HOD, bullock_02,berlind_02,bosch03,berlind_weinber_03,kravtsov_04,zehavi_weinberg_04, hamana_04,zehavi_zheng_05,zheng_05,hamana_06,conroy_wechsler_06,lee_09}. In our work we calculate the HOD without assuming any parametric form. Combining this with the dark matter halo abundance, bias and density profile, the galaxy clustering can be calculated.  To begin with we assume that each halo can host at most one visible central galaxy. By fitting the observed LF of LBGs, we find the masses of dark matter halos which host an LBG of a given luminosity. We then show that our prescription of star formation that fits the observed LF of LBGs can also simultaneously explain their large scale clustering ($\\theta \\ge 80\"$). In order to also account for the small angular scale clustering we calculate how many subhalos hosting a detectable satellite can form in a bigger parent halo using the conditional mass function \\citep{cooray_sheth_02} and our star formation prescription. Thus we provide a method of calculating the HOD from first principles, which can then be used to predict the LBG clustering. Using our approach, we show that one can explain both the UV LFs and luminosity dependent clustering of LBGs and gain useful insights into galaxy formation.  The organization of this paper is as follows. In the next section we describe our physical model for computation of the LF of LBGs. In Section 3 we focus on the clustering of LBGs on large angular scales. We then turn to our physically motivated model to calculate the central and satellite contributions to the HOD and use these to obtain total angular correlation functions at all angular scales. Section 5 presents a comprehensive comparison of the total angular correlation function computed in various models with observations. A discussion of our results and conclusions are presented in the final section. For all calculations we adopt a flat $\\Lambda$CDM universe with cosmological parameters consistent with 7 year Wilkinson Microwave Anisotropy Probe (WMAP7) observations \\citep{larson_11_wmap7}. Accordingly we assume $\\Omega_m=0.27$, $\\Omega_\\Lambda=0.73$, $\\Omega_b=0.045$, $h=0.71$, $n_s = 0.963$ and $\\sigma_8 = 0.801 h^{-1}$Mpc. Here $\\Omega_i$ is the background density of any species 'i' in units of critical density $\\rho_{c}$. The Hubble constant is $H_0 = 100 h$ km s$^{-1}$ Mpc$^{-1}$ ",
        "conclusions": "We have presented here a physically motivated semi-analytical model of galaxies to understand the clustering of high redshift Lyman break galaxies, where the model parameters are constrained by the observed high z luminosity function. For this purpose we use and expand upon the standard halo model. Galaxies are assumed to be formed inside dark matter halos. Their luminosity is determined by a physical model of star formation, which is a function of the mass and age of the hosting halo.  We began by assuming that each halo can host at most one visible galaxy. On fitting the observed LF, we determine the relationship between the luminosity of a galaxy and the mass of its host halo. This allows us to calculate the large scale bias for LBGs satisfying any luminosity threshold. This bias is then folded in with the dark matter power spectrum to predict the two point spatial correlation function of these LBGs and hence the angular correlation function $w(\\theta)$. For our fiducial model, we find that the predicted $w(\\theta)$ compares very well with the observed data of both \\citet{hildebrandt_09_acf} and \\citet{ouchi_hamana_05_acf} at $\\theta > 80''$, for the whole range of redshifts $z=3-5$ and limiting luminosities (magnitudes) (see Figs.~\\ref{fig:acfz_lin},\\ref{fig:acfz_full}, \\ref{fig:acfz4_ouchi} and Table~\\ref{tab3}).  The predicted large scale galaxy bias $b_g$ agrees well with that observationally determined by \\citet{hildebrandt_09_acf} (see Table~\\ref{tab1}). At a given $z$ the bias increases with increasing luminosity, while for a given $L_{th}$ it increases with $z$. However, we find a smaller spread in $b_g$ as a function of $L_{th}$ at any given $z$ compared to that of \\cite{hildebrandt_09_acf}. Remarkably, we find that the predicted large scale clustering of LBGs is fairly insensitive to the assumed astrophysical or cosmological parameters, provided we simultaneously fit the observed LF. This may point to an important internal consistency of our physical model; that if we fix the mass to light ratio correctly by using the LFs of LBGs, then the standard LCDM model correctly predicts the amplitude of their large scale clustering.  We then extended our approach by incorporating the Halo occupation distribution, which provides the distribution of galaxies inside dark matter halos. This is separated into the central, $f_{cen}$, and satellite, $N_s$, contributions. Often the HOD is modelled in a parametrized form. Instead we have adopted a more physical approach. We use our prescription for computing the LF to estimate $f_{cen}$. The conditional mass function and our star formation model are used to calculate the mean number and luminosity of satellite galaxies in a parent halo. An additional parameter, $\\Delta t_0$, is introduced in calculating $N_s$ (see Table~\\ref{tab0}). If parent halo collapses at a time $t_p$, a subhalo can host a detectable galaxy only if it collapses at an earlier epoch, $t_s < t_p - \\Delta t_0$. In order to explain the small angular scale clustering, one requires $\\Delta t_0$ of order $t_{dyn}$, the dynamical time scale of the parent halo. A much smaller (or larger) $\\Delta t_0$ leads generically to an excess (deficit) of small angular scale clustering.  The calculated forms of $f_{cen}(M,L_{th},z)$ and $N_s(M,L_{th},z)$ compare reasonably with that assumed in parametrized models of \\citet{hamana_06} and the simulations of \\citet{conroy_wechsler_06}. The average value of $f_{cen}$ is typically $0.4$. Thus, on the average, $40\\%$ of the halos above a minimum mass $M_{min}$ (which itself depends on the luminosity threshold), at any given redshift, can host detectable central galaxies. Further the average value of $N_s$ is about $0.02-0.04$, or about $5-10\\%$ of the halos with detectable central galaxies, also will have a detectable satellite. Indeed it is such pairs which contribute to the small angular scale clustering. At $z = 4$ and for apparent magnitude thresholds in the range 25-26, the average mass of halos contributing to the observed clustering, $M_{av}$, ranges from $10^{12} M_\\odot -3.9 \\times  10^{11} M_\\odot$. At $z = 5$ these masses are smaller by a factor $\\sim 1.4$, while at $z = 3$ these masses are larger by a similar factor. At any given redshift and magnitude threshold, the typical mass $M_P$ of parent halos hosting detectable satellite galaxies, are about 2 times larger than $M_{av}$.  Having obtained the HOD, we can calculate both the one halo and two halo contributions to the total correlation function of LBGs. Our simple physical model gives a reasonable fit to the observed clustering of LBGs at all angular scales for the faintest LBGs with $m \\ge 25$ at $z=3$ and for $m \\ge 26.5$ at $z=4$. At $z=5$, and for the most luminous galaxies at $z=3,4$, the predicted $w(\\theta)$ again fits the observed data well at both large ($\\theta>80''$) and small ($\\theta < 10''$) angular scales. The clustering at small angular scales as mentioned above, is likely to be dominated by pairs of LBGs rather than rich clusters, as the number of detectable galaxies hosted by most collapsed halos is typically less than 2.  The small angular scale clustering is also not very sensitive to changes of several cosmological and astrophysical parameters from their fiducial values, as long as we simultaneously fit the observed LF. However, the amplitude of $w(\\theta)$ on small angular scales is very sensitive to the value of $\\kappa$, which determines the duration of star formation activity in a halo. The present data are consistent with $\\kappa \\sim 1$ or a star formation duration of the order of the dynamical time scale $t_{dyn}$ of the dark matter halo.  We find that the following broad physical picture of LBGs consistently accounts for their observed LF and clustering. First the average mass of the halos hosting the brightest central LBGs at $z=3-5$, with $-21< M_{AB} < -20$, is around $3\\times 10^{11} M_\\odot$ to $1.5 \\times 10^{12} M_\\odot$. Halos which host detectable satellites and contribute dominantly to the small angular clustering are more massive by a factor of 2 or so. Typically fainter LBGs or those at higher $z$ are hosted by smaller mass halos. In these galaxies about $50\\%$ of the stars are formed over a timescale of $300-500$ Myr for $z=5-3$, by converting $\\sim 3-8\\%$ of the baryons to stars. Our physical model for the HOD suggests that approximately $40\\%$ of the halos above a minimum mass $M_{min}$, can host detectable central galaxies. This is comparable to the duty cycle values preferred by \\citet{lee_09,lee_ferguson_12}. Further, about $5-10\\%$ of these halos are likely to also host a detectable satellite. These satellites form over a dynamical timescale or so prior to the formation of the parent halo. The small angular scale clustering is mainly due to central-satellite pairs. The average fraction of halos which can host a central LBG can be compared to the duty cycle invoked in the literature. Finally, a preliminary study suggests that the star formation model that we have invoked is also consistent with observation of the SFR-$M_*$ relation, and the stellar mass function.   The theoretically predicted $w(\\theta)$ at intermediate angular scales is smaller than that observed, for the brightest LBGs at $z=3,4$ and at $z=5$. We explored in detail whether this excess can be due to a more extended satellite galaxy distribution. This only partly accounts for the discrepancy. We also find that the  non-linear corrections to dark matter power spectrum does not raise the predicted clustering at the intermediate scales to the observed values. Note that the typical mass of dark matter halo hosting a galaxy increases with its luminosity. Also the higher order (quasi-linear) corrections to the dark matter halo bias is expected to be larger for higher mass and higher redshift halos, which are rarer. Therefore we suspect that the higher order quasi-linear corrections to galaxy bias could be playing a role in explaining the excess intermediate scale clustering, a possibility which we hope to explore in the future. Nevertheless, it is noteworthy that the predictions from our simple physical model, employing only a few free parameters, can fit the observed clustering data over a wide range of scales, redshifts and limiting luminosities."
    },
    "1208/1208.5481_arXiv.txt": {
        "abstract": "Weakly interacting massive particles (WIMPs) remain a prime candidate for the cosmological dark matter (DM), even in the absence of current collider signals that would unambiguously point to  new physics below the TeV scale. The self-annihilation of these particles in astronomical targets may leave observable imprints in cosmic rays of various kinds. In this review, we focus on gamma rays which we argue to play a pronounced role among the various possible messengers. We discuss the most promising spectral and spatial signatures to look for, give an update on the current state of gamma-ray searches for DM and an outlook concerning future prospects. We also assess in some detail the implications of a potential signal identification for particle DM models as well as for our understanding of structure formation.  Special emphasis is put on the possible evidence for a 130 GeV line-like signal that we recently identified in the data of the Fermi gamma-ray space telescope. ",
        "introduction": "Evidence for a sizable non-baryonic and cold dark matter  (DM) component  in the universe derives from an impressive range of unrelated cosmological observations \\cite{1208.3662}, covering distance scales from tens of kpc to several Gpc and leaving very little room for alternative explanations. On cosmological scales, DM  contributes a fraction of $\\Omega_\\chi = 0.229 \\pm 0.015$ to the total energy density of the universe \\cite{Komatsu:2010fb}.  Weakly Interacting Massive Particles (WIMPs) provide a theoretically particularly appealing class of candidates for the so far obscure nature of DM \\cite{Jungman:1995df,Bergstrom:2000pn, Bertone:2004pz}, with the lightest supersymmetric neutralino often taken as a useful template for such a WIMP.  It is often argued that the thermal production of WIMPs in the early universe generically leads to a relic density that coincides with the observed order of magnitude of $\\Omega_\\chi$, though this rests on the assumption of a standard cosmological expansion history and  there exist  well-motivated particle physics scenarios that predict alternative production mechanisms for WIMP DM \\cite{Feng:2010gw}.  While the LHC non-observation of new particles below the TeV scale (apart, possibly, from the Higgs boson) has already prompted  doubts whether the WIMP DM scenario is still our best bet \\cite{Bertone:2010at}, it must be stressed that electroweak low-energy observables ($g-2$ in particular) do  favor new physics contributions not too far above 100\\,GeV \\cite{Czarnecki:2001pv}. While this tension starts to considerably disfavor very constrained models of, e.g., supersymmetry \\cite{Bechtle:2012zk}, it may simply be an indication that the new physics sector which the WIMP belongs to appears at a much smaller mass scale than any new colored sector.  Attempts to identify WIMP DM can be classified into collider searches for missing transverse energy,  direct searches for the recoil  of WIMPs off the nuclei of terrestrial  detectors and indirect  methods that aim at spotting the products of WIMP self-annihilation. Among possible messengers for such indirect searches, \\emph{gamma rays} play a pronounced role as they  propagate essentially unperturbed through the galaxy and therefore directly point to their sources, leading to distinctive \\emph{spatial signatures}; an even more important aspect, as we will see, is the appearance of pronounced \\emph{spectral signatures}.  This prime role of gamma rays provides our motivation for an updated and dedicated review on these messengers, which we hope will prove useful and complementary to existing general reviews on indirect DM searches \\cite{Cirelli:2010xx,Profumo:2010ya,Lavalle:2012ef}. Indeed, the recent indication for a DM signature in gamma-ray observations of the Galactic center (GC) \\cite{Bringmann:2012vr, Weniger:2012tx} makes such a review extremely timely, and we therefore dedicate a considerable part of it to discuss in great detail both the status of the potential signal and its implications.  Gamma rays can either be observed directly from space or, via the showers of secondary particles they trigger in the atmosphere, indirectly with ground-based experiments. The former option necessarily implies rather small effective areas and  an upper bound on the photon energy that can reliably be resolved, but allows for a large field of view and the observation of gamma rays at comparably small energies. Particularly promising instruments for the latter option are imaging Air Cherenkov Telescopes (IACTs) that detect the Cherenkov light emitted by the shower particles and use efficient image reconstruction algorithms to determine the characteristics of the primary photon. These instruments have a limited field of view and a lower energy threshold set by the need of discriminating photons from the background of primary muons and hadronic cosmic rays; their extremely large effective area and rather small field of view make them ideal for pointed observations. In Table \\ref{tab:telescopes}, we pick typical examples for space- and ground-based experiments that are currently operating or planned for the future and compare some basic telescope characteristics that are particularly relevant for DM searches.  Experiments that fall into the same broad categories but are not listed explicitly in the Table include for example AGILE~\\cite{Tavani:2008sp} and VERITAS \\cite{Holder:2006gi}, as well as the future CALET \\cite{CALET}  and DAMPE~\\cite{DAMPE, Li:2012qg}. We stress that the numbers in Tab.~\\ref{tab:telescopes} are intended to provide a convenient order-of-magnitude comparison of instrumental characteristics; they should \\emph{not} be used as the basis of detailed sensitivity estimates (see, however, the stated references).  \\begin{table}[t!] \\scriptsize \\centering \\begin{tabular}{llcccccc} \\toprule   & Time of   & $E$-range & $A_\\text{eff}$ & \\!\\!\\!\\!\\!\\!Sens.\\!\\!\\! & $\\Delta E/E$ & F.O.V. & $\\Delta \\theta$  \\\\ & operation & [GeV]        & [m$^2$]        &  \\!\\!\\!\\!\\!\\! [10$^8$m$^{2}$s]$^{\\text{-}1}$\\!\\!\\!     &   [\\%]      & [sr]   & [$^\\circ$]  \\\\ \\midrule Fermi-LAT   & 2008--$2018^*$  & 0.2--300 & 0.8 & 200 & 11 & 2.4  & 0.2         \\\\ AMS-02/Ecal & 2011--$2021^*$  & 10--1000    & 0.2  & 1000  & 3  & 0.4  & 1.0         \\\\ AMS-02/Trk  & 2011--$2021^*$  & 1--300      & 0.06 & 1000  & 15 & 1.5  & 0.02        \\\\ GAMMA-400   & $2018^*$--\\dots & 0.1--3000   & 0.4  & 100 & 1  & 1.2  & 0.02 (0.006)       \\\\ \\midrule MAGIC   & 2009--\\dots    & $\\gtrsim50$ & $2\\!\\cdot\\!10^4(7\\!\\cdot\\!10^4)$ & $10(0.2)$    & 20(16)  & 0.003      & 0.17(0.08)  \\\\ HESS-II & 2012--\\dots    & $\\gtrsim30$ & $4\\!\\cdot\\!10^3(10^5$)       & 4(0.1) & 15(15) & 0.003     & 0.13(0.07) \\\\ CTA     & $2018^*$--\\dots & $\\gtrsim20$ & $5\\!\\cdot\\!10^4(10^6)$       & 1(0.02) & 20(10)   & $>0.006$ & 0.1(0.06)  \\\\ \\bottomrule &&&&&&& \\hfill $^*$ planned \\end{tabular} \\caption{\\label{tab:telescopes}  Rough comparison of basic telescope characteristics relevant for indirect DM searches with gamma rays, for a selection of typical space- and ground-based experiments that are currently operating, shortly upcoming or planned for the future. The quoted sensitivity is for point sources at the $5\\sigma$ level, after 1yr  (50 hrs) of space- (ground-) based observations and assuming typical backgrounds. Where applicable, numbers refer to photon energies at or above $E \\simeq 100\\GeV$ ($1\\TeV$). The angular resolution $\\Delta\\theta$ denotes the $68\\%$ containment radius. More details in Refs.~\\cite{Fermi:performancePASS7} (Fermi-LAT), \\cite{Jacholkowska:2005nz,Battiston:1999yb,AMS:web} (AMS-02), \\cite{Galper:2012fp,Topchiev,Galper:2012ji} (GAMMA-400), \\cite{Aleksic:2011bx} (MAGIC), \\cite{Hess2:Hd2012} (HESS-II) and \\cite{CTA} (CTA).} \\end{table}  The expected DM-induced gamma-ray flux from a direction $\\psi$, averaged over the opening angle $\\Delta\\psi$ of the detector, is given by \\be \\label{flux} \\frac{d\\Phi_{\\gamma}}{dE_\\gamma} (E_\\gamma,\\psi) = \\frac{1}{8\\pi} {\\int_{\\Delta\\psi}\\frac{d\\Omega}{\\Delta\\psi}\\int_\\mathrm{l.o.s} d\\ell(\\psi) \\rho_\\chi^2(\\mathbf{r})} \\times \\left({\\frac{\\langle\\sigma v\\rangle_\\mathrm{ann}}{m_{\\chi}^2} \\sum_f B_f\\frac{dN_\\gamma^{f}}{dE_\\gamma}}\\right) \\,, \\ee where the integration is performed along the line of sight (l.o.s.), $\\langle\\sigma v\\rangle_\\mathrm{ann}$ is the average velocity-weighted annihilation cross section, $m_\\chi$ the mass of the DM particle (for which we assume $\\chi=\\bar\\chi$), $\\rho_\\chi$ the DM density, $B_f$ the branching ratio into channel $f$ and  $N_\\gamma^{f}$ the number of photons per annihilation. An often quoted reference value for $\\langle\\sigma v\\rangle_\\mathrm{ann}$ is the so-called `thermal cross section' of $\\langle \\sigma v\\rangle\\sim3\\cdot10^{-26}{\\rm cm}^3{\\rm s}^{-1}$,  which is the annihilation rate expected for thermally produced WIMPs in the most simple case (i.e.~$s$-wave annihilation without resonances or co-annihilations \\cite{Griest:1990kh}). The right part (in parentheses) of Eq.~(\\ref{flux})  contains all the particle physics input and, for the typically very small DM velocities, is usually sufficiently independent of $v(\\mathbf{r})$ that it can be pulled outside the integrals (note, however, that this is \\emph{not} true for strongly velocity-dependent cross-sections like in the case of Sommerfeld enhancement \\cite{sommerfeld,Hisano:2003ec,Hisano:2004ds,ArkaniHamed:2008qn}, resonances or $p$-wave annihilation). It contains the full \\emph{spectral} information that we will discuss in some detail in Section \\ref{sec:spec}. The remaining part, sometimes referred to as the astrophysical factor (or '$J$-value', with $J\\equiv\\int d\\Omega\\int\\!\\!d\\ell\\,\\rho_\\chi^2$), contains in that case the full information about the \\emph{spatial} distribution of the signal and will be discussed  in Section \\ref{sec:spatial}.  We  continue by reviewing in Section \\ref{sec:status}  gamma-ray limits on DM annihilation as well as  the current status of claimed DM signals. The potentially enormous  implications of a signal identification  for our understanding of both the underlying particle model and structure formation are then outlined in Section \\ref{sec:implications}, with a focus on  the intriguing 130\\,GeV feature in the direction of the GC. We discuss future prospects for  the detection of DM with gamma rays in Section \\ref{sec:prospects} and conclude in Section \\ref{sec:conc}. For most of this review, we will assume that DM consists of WIMPs; many aspects, however, can be applied -- or generalized in a straight-forward way -- to other cases as well, most notably decaying DM \\cite{Buchmuller:2007ui,Cirelli:2012ut} (for which one simply has to replace $\\frac12\\langle\\sigma v\\rangle\\rho_\\chi^2\\rightarrow m_\\chi\\Gamma\\rho_\\chi$ in Eq.~(\\ref{flux}), where $\\Gamma$ is the decay rate).  Where applicable, we will comment on this on the way. ",
        "conclusions": "\\label{sec:conc}  In this review, we have argued that one may consider gamma rays as the \\emph{golden channel} of indirect searches for DM in view of the extraordinarily rich spectral and angular information they can carry. This does not only help to discriminate signals from backgrounds but could eventually reveal valuable details about the properties of the DM particles. We have discussed the most important signatures in quite some detail and provided an update on current limits, demonstrating that indirect searches start to probe realistic cross sections and thus become competitive probes of physics beyond the standard model.  While too early for a final judgement at the time of this writing, the line feature at 130\\,GeV that is seen in the Fermi data might turn out to be the most promising DM \\emph{signal} claimed so far. In fact, the intrinsic width of this feature must be smaller than  roughly 20\\% (18\\% at 95\\%CL) -- which leaves  lines, VIB or box signals (Fig.~\\ref{fig:spec_comp}) as possible channels for an explanation in terms of DM. On the other hand, it is extremely challenging to find any explanation related to astrophysics for such a spectral feature; for example, even a very hard contribution to the gamma-ray flux, with a sharp break at 130\\,GeV, cannot describe the data in a satisfactory way (Fig.~\\ref{fig:powerlaw}). The signal morphology is perfectly consistent with annihilating DM and an Einasto profile for the DM density, at least for distances larger than the possible displacement from the GC by 1-2 degrees, but essentially rules out both cored and more contracted profiles (Fig.~\\ref{fig:profile}); decaying DM is also in strong tension with the data. The data  show a weak hint for a second peak at 114\\,GeV (Fig.~\\ref{fig:gggZ} and Tab.~\\ref{tab:line_limits}) which is exactly the combination of energies expected for the annihilation of 130\\,GeV DM particles into $\\gamma\\gamma$ and $\\gamma Z$ final states. However, large annihilation rates into these channels rather generically imply large annihilation rates rate also at tree-level, in potential conflict with continuum gamma-ray limits; VIB, on the other hand, does not suffer from this drawback (Fig.~\\ref{fig:SUSY_scan}).  If confirmed by the Fermi collaboration or other experiments, and in the absence of satisfactory instrumental or astrophysical explanations, this signal would lead to the exciting conclusion that the first particle beyond the standard model has been found in space rather than at a collider. We have discussed at length how astrophysical observations would already now help to determine detailed properties of this new particle. The situation will further improve  in the relatively near future given that  prospects to study the 130\\,GeV  feature in more detail are extremely good. However, we believe that even if the DM origin of the signal is eventually not confirmed, our analysis serves to make a compelling case for the importance of focussing on clear spectral features in future searches for DM.  During the last ten years or so, most limits on DM annihilation have improved by about one order of magnitude and this trend is expected to continue for the \\emph{next decade} (Fig.~\\ref{fig:time}). We have further estimated the systematics-limited (or `fundamental') reach of gamma-ray experiments with present technology, demonstrating that there is still quite some room for improvement even beyond those limits expected for the next decade, especially for space-based instruments (but also for ground-based telescopes \\emph{if} the cosmic-ray electron background  can at least partially be rejected). Eventually, it may thus in principle be possible to probe cross sections down to at least one order of magnitude below the thermal value for TeV-scale particles; for many models, this would correspond to interactions too feeble  to show up in any other kind of experiment, including direct or collider searches. While even those limits may not be sufficient to completely close the window for WIMP DM, model-building would certainly need to become increasingly sophisticated to avoid them.  Let us finally stress that in order to fully identify the properties of the DM particles, it will of course be indispensable to correlate a suspected DM signal  in gamma rays with results from indirect searches at other wavelengths and with other messengers. The same holds for direct searches and new data from colliders, both of which are guaranteed to deliver substantial new input in the near future -- be it in terms of greatly improved limits or actual first hints for a signal. Chances are thus high that the next decade will either bring us a great deal closer to the long-sought nature of DM or, in the most pessimistic scenario in terms of detectional prospects, force us to seriously question the very idea of DM being composed of WIMPs.  \\bigskip"
    },
    "1208/1208.1094_arXiv.txt": {
        "abstract": "Chiral effective field theory (EFT) provides a systematic expansion for the coupling of WIMPs to nucleons at the momentum transfers relevant to direct cold dark matter detection. We derive the currents for spin-dependent WIMP scattering off nuclei at the one-body level and include the leading long-range two-body currents, which are predicted in chiral EFT. As an application, we calculate the structure factor for spin-dependent WIMP scattering off $^{129,131}$Xe nuclei, using nuclear interactions that have been developed to study nuclear structure and double-beta decays in this region. We provide theoretical error bands due to the nuclear uncertainties of WIMP currents in nuclei. ",
        "introduction": "Cosmological and astrophysical observations have established that more than $20 \\%$ of the energy density of our universe is dark matter, a rarely interacting nonbaryonic form of matter, whose specific composition remains unknown~\\cite{theo}. Promising candidates are weakly interacting massive particles (WIMPs), such as neutralinos, the lightest supersymmetric particles predicted by extensions of the standard model. This has spurred direct detection of cold dark matter via elastic scattering off nuclei, requiring detailed knowledge of the response to WIMP induced currents in nuclei. This presents a challenging problem, because even if the coupling of neutralinos (or other particles) to quarks is known, it needs to be evaluated at the nucleus level in the nonperturbative regime of quantum chromodynamics~\\cite{new}.  Chiral effective field theory (EFT) provides a systematic expansion for nuclear forces and the coupling to external probes for momenta of order of the pion mass, $p \\sim m_{\\pi} \\sim 100 \\mev$, which are typical momentum transfers in direct dark matter detection. In this paper, we focus on spin-dependent neutralino scattering off nuclei, which is used to constrain WIMP properties~\\cite{exp} and is particularly sensitive to nuclear structure. We derive the currents for spin-dependent WIMP scattering at the one-body level and include the leading long-range two-nucleon currents, which are predicted in chiral EFT. Similar weak neutral currents are key for providing accurate predictions of neutrino breakup of the deuteron for SNO~\\cite{SNO}, while the corresponding weak charged currents have been found relevant in Gamow-Teller and double-beta decays of medium-mass nuclei~\\cite{Javier}.  We apply the developed chiral EFT currents to calculate the structure factor for spin-dependent WIMP scattering off $^{129,131}$Xe, as xenon isotopes provide the tightest limits on WIMP couplings~\\cite{XENON}. To describe these nuclei, we employ the largest many-body spaces accessible with nuclear interactions previously used to study nuclear structure and double-beta decays in this region~\\cite{Menendez}. ",
        "conclusions": "This presents the first calculation of spin-dependent WIMP currents in nuclei based on chiral EFT, including the leading long-range 2b currents. They predict a $25-55 \\%$ reduction of the isovector part of the one-body axial-vector WIMP currents, where the range provides an estimate of the theoretical uncertainties of WIMP currents in nuclei. This should be included in limits on the WIMP couplings, where the spin-dependent analysis provides complementary constraints~\\cite{exp}. As an application, we have calculated the structure factors for spin-dependent WIMP scattering off $^{129,131}$Xe nuclei, using the largest valence spaces accessible with nuclear interactions that have been tested in nuclear structure and double-beta decay studies in this region. Future work includes developing consistent interactions based on chiral EFT, where the present frontier is in the calcium region~\\cite{Gallant}, and investigating other nuclei and responses."
    },
    "1208/1208.1080_arXiv.txt": {
        "abstract": "We have conducted $VI$ CCD photometry of the two open clusters NGC 1245 and NGC 2506 using the CFH12K CCD camera. Our photometry covers a sky area of $84^\\prime \\times 82^\\prime$ and $42^\\prime \\times 81^\\prime$ for the two clusters, respectively, and reaches down to $V \\approx 23$. We derived the physical parameters using detailed theoretical isochrone fittings using $\\chi^2$ minimization. The derived cluster parameters are $E(B-V) = 0.24 \\pm0.05$ and $ 0.03 \\pm0.04$, $(V-M_V)_0  = 12.25 \\pm0.12$ and $12.47 \\pm0.08$, $age(Gyr) = 1.08 \\pm0.09$ and $2.31 \\pm0.16$, and [Fe/H]= $-0.08 \\pm0.06$ and $-0.24 \\pm0.06$, respectively for NGC 1245 and NGC 2506, We present the luminosity functions (LFs) of the two clusters, which reach down to $M_V \\approx 10$, and derive mass functions (MFs) with slopes of $ \\Gamma = -1.29$ for NGC 1245 and $ \\Gamma = -1.26$ for NGC 2506. The slopes are slightly shallower than that of the solar neighbourhood, implying the existence of dynamical evolution that drives the evaporation of the low-mass stars in the clusters. ",
        "introduction": "%  Open clusters are located in the Galactic disc as Population  I, and it is important to investigate their properties and the spatial distribution of the Galaxy to understand the Galaxy structure and evolution, especially the formation and dynamical evolution of the Galactic disc. Because stars in the Galactic disc are originated in evaporated stars from open clusters, the luminosity function of open clusters is an important parameter for understanding their formation and evolution. We can easily obtain the LF from CCD observations, but an analysis of the LF is not easy because the observed LF of open clusters reflects the dynamical evolution which depends on the cluster's age as well as the initial mass function (IMF).  It is necessary to have a deep photometry over the entire cluster field to understand the LF of old open clusters down to the magnitudes of the evaporating stars because they are much fainter than the turn-off stars and are supposed to be located around the tidal radius of the clusters. But there are only a few open clusters of which photometry is deep and wide enough to study the effect of dynamical evolution on the structure of open clusters. Thus, we selected open clusters which are old enough for the low mass stars to evaporate from the clusters by dynamical evolution and rich enough suitable for a detailed analysis of the spatial distribution of low mass stars. We also took into account the distance of open clusters to ensure that our photometry reaches up to $M_V \\sim 10$.  The old open cluster NGC 1245 is located in the Perseus arm at a Galactic latitude of $-8^{\\circ}.93$. Using photographic photometry data from the Schmidt telescope, \\citet{per78} discussed the radial density distribution of stars with a limiting magnitude of $B = 17$ and a radius of $r \\approx 83^{\\prime}$. They presented cluster boundaries with a nucleus ($r < 13^{\\prime}$) and corona ($ 13^{\\prime} < r < 61^{\\prime}$). \\citet{car94} and \\citet{sub03} found a shallower slope of the mass function of NGC 1245 than \\citet{sal55}. Using $BV$ CCD photometry, \\citet{sub03} presented the surface number density with a limiting magnitude of $V = 17$ and a radius of $r \\sim 6^{\\prime}.5$. They also indicated a lack of stars in the cluster centre. \\citet{bur04} derived physical parameters and a core radius of $R_c = 3^{\\prime}.1$ from $BVI$ CCD photometry. They determined an optical absorption of $A_V = 0.68$ magnitude and a distance modulus $(V-M_V)_0 = 12.27$; they also reported that the cluster has no differential reddening, in contrast to previous studies such as \\citet{car94} and \\citet{wee96}.  The old open cluster NGC 2506 also has rich members located at a Galactic latitude of $+9^{\\circ}.94$. A radial density profile of the proper motion membership stars obtained by \\citet{chi81} indicated that bright stars are more concentrated than faint stars in this cluster. \\citet{mar97} suggested that more than 20\\% of the main sequence stars are binary. From a photometric study by \\citet{hen07}, the distance of this cluster is known to be 3.4 kpc and its age is 1.79 Gyr.  This paper presents the deep $VI$ CCD photometry of the two old open clusters NGC 1245 and NGC 2506, covering the entire cluster fields including the surrounding field regions for field star correction, along with the derivation of the cluster parameters such as age, metallicity, distance, and interstellar reddening. We will discuss the dynamical structures of the open clusters in a forthcoming paper (Lee, Kang $\\&$ Ann, in preparation), but here we present the physical properties, luminosity functions, and mass functions of the two open clusters along with their physical properties.  The present paper is organized as follows. In \\S2, we describe the observations and data reductions. We present the physical parameters of the clusters in \\S3, and describe the luminosity functions and mass functions in \\S4 and \\S5, respectively. A summary is given in the final section.  \\begin{figure} \\vspace{2mm} \\epsfig{figure=fig02_shlee.eps, height=0.96\\textwidth, width=0.48\\textwidth} \\vspace{2mm} \\caption{ Observed fields ($42^{\\prime} \\times 81^{\\prime}$) of NGC 2506. We combined 3 V-filter images, where one CCD image has a size of $42^{\\prime} \\times 28^{\\prime}$, to make the mosaic image. North is up and east is to the left. } \\vspace{-2mm} \\label{fig_fov3} \\end{figure} ",
        "conclusions": "We  conducted deep and wide $VI$ photometry of old open clusters NGC 1245 and NGC 2506 based on CCD observations using the CFHT. The present photometry is suitable for investigating dynamical evolution in open clusters because it is deep enough to reach down to $M_{V}\\approx10$ and wide enough to cover the entire regions of the two clusters including the field regions for an effective correction of the field star contamination.  We derived the physical parameters of the two clusters using detailed isochrone fittings based on $\\chi^2$ minimization: $E(B-V) = 0.24 \\pm0.05$ and $ 0.03 \\pm0.04$, $(V-M_V)_0  = 12.25 \\pm0.12$ and $12.47 \\pm0.08$, $age(Gyr) = 1.08 \\pm0.09$ and $2.31 \\pm0.16$, and [Fe/H]= $-0.08 \\pm0.06$ and $-0.24 \\pm0.06$ for NGC 1245 and NGC 2506, respectively. The quoted errors are the total errors which are quadratic sum of statistical errors associated with the $\\chi^2_{min}$ and systematic errors related to the assumed parameters and theoretical models. The present estimates of the physical parameters of two clusters are in good agreement with the previous estimates.  We also derived the LFs of NGC 1245 and NGC 2506. The LF of NGC 1245 shows a flat profile between $M_V=3$ and $M_V=10$, whereas the LF of NGC 2506 displays a slight rise. This difference seems to be not due to the errors related to the field star correction but due to the dynamical structures, which can differ because dynamical evolution can be different for clusters of a similar age as it is affected by the environment as well as by the internal properties of the cluster. We derived the mass functions of NGC 1245 and NGC 2506, which show slightly shallower slopes than that of the solar neighbourhood IMF. The shallower slope is understandable if we consider that the derived mass function is the present day mass function, which is thought to be different from the IMF of the clusters due to dynamical evolution, since the ages of NGC 1245 and NGC 2506 are old enough for mass segregation and evaporation of low-mass stars. We will discuss the dynamical properties, structures and halo of these two clusters in a forthcoming paper."
    },
    "1208/1208.0423_arXiv.txt": {
        "abstract": "{We investigated the environment of the infrared dust bubble N68 and searched for evidence of triggered star formation in its surroundings. We performed a multiwavelength study of the nebula with data taken from several large-scale surveys: GLIMPSE, MIPSGAL, IRAS, NVSS, GRS, and JCMT. We analyzed the spectral profile and the distribution of the molecular gas ($^{13}$CO J = 1 - 0 and J = 3 - 2), and the dust in the environment of the N68. The position-velocity diagram clearly shows that the N68 may be expanding outward. We used two three-color images of the mid-infrared emission to explore the physical environment, and one color-color diagram to investigate the distribution of young stellar objects (YSOs). We found that the 24 $\\mu$m emission is surrounded by the 8.0 $\\mu$m emission. Morphologically, the 1.4 GHz continuum correlates strongly with the 24 $\\mu$m emission, and the $^{13}$CO J = 1 - 0 and J = 3 - 2 emissions correlate well with the 8.0 $\\mu$m emission. We investigated two compact cores located at the shell of the N68. The spectral intensity ratios of $^{13}$CO J = 3 - 2 to J = 1 - 0 range from 5 to 0.3. In addition, young star objects, masers, IRAS, and UC HII regions distribute at the shell of bubble. The active region may be triggered by the expanding of the bubble N68. ",
        "introduction": "% \\label{sect:intro}   There are many signatures of star formation, for example, outflow/inflow, dark cloud and HII region, which can be used to investigate the process of star formation. Also there are many different kinds of interactions which can trigger the star formation, such as cloud-cloud collision\\citep{li2012}, supernova explosion \\citep{xujl2011,xu2011}, bubble expansion \\citep{s51} and so on. \\citet{chur2006,chur2007} detected and cataloged about 600 mid-infrared dust (MIR) bubbles between longitudes -60$^{\\circ}$ and +60$^{\\circ}$. The bubbles have bright 8.0 $\\mu$m shells that enclose bright 24 $\\mu$m interiors. The infrared (IR) dust bubbles may be produced by exciting O- and/or B-type stars, which is located inside bubble. The ultraviolet (UV) radiation from exciting stars may heat dust and ionize the gas to form an expanding bubble shell \\citep{wats2008}, which is known as ''collect-and-collapse'' process. This process can trigger a massive star formation near the shell clumps. Therefore, the bubbles present an important opportunity to study the interaction between the HII regions and molecular clouds.   A few individual bubbles have been studied well, such as N49 and S51. There are many models and observations to explain the dusty wind-blown bubbles, such as bubble N49 of \\citet{ever2010}. Recently, we reported an expanding bubble S51 shown with a shell and a front side, employing $^{13}$CO and C$^{18}$O J = 1 - 0 emission lines \\citep{s51}. \\citet{beau2010} reported CO J = 3 - 2 maps of 43 Spitzer identified bubbles. \\citet{wats2008} present an analysis of wind-blown, parsec-sized, mid-infrared bubbles and associated star formation.   To add the investigated bubble examples, we selected the IR dust bubble [CPA2006] N68 (hereafter N68) from the catalog of \\citet{chur2006}. N68 is a complete or closed ring and centered on Galactic $l$=35.654, $b$=-0.062 (or R.A.(J2000) = 18$^{h}$56$^{m}$25$^{s}$.70, DEC.(J2000) = 02$^{\\circ}$26$'$01$''$.0). It has a distance of 10.6 kpc, which was obtained by \\citet{ande2009} based on comparing the velocity of the ionized gas with the maximum velocity of HI absorption, and one looking for HI absorption at the velocity of molecular emission. In addition, its size is 34 pc $\\times$ 17 pc and the eccentricity of the ellipse is 0.72.   In this work, we mainly present a multiwavelength study of the environment surrounding the IR dust bubble N68. We explore its surrounding interstellar medium (ISM) and search for signatures of star formation. The observations and data are described in Sect.\\ref{sect:data}; the results and discussions about the bubble N68 environment are presented in Sect.\\ref{sect:results}; Sect.\\ref{sect:summary} summarizes the results. ",
        "conclusions": "\\label{sect:summary}  We have investigated the environment of the IR dust bubble N68 with the $^{13}$CO J = 1 - 0 and J = 3 - 2 emission lines, the 1.4 GHz continuum, and the IR bands emission. The main results can be summarized as follows.  The morphology of the molecular line emission ($^{13}$CO J = 1 - 0 and J = 3 - 2) and the associated velocity signatures are consistent with the shell structure seen from the 8.0 $\\mu$m images. The expanding shells in N68 are suggestive of triggered massive star formation. There is evidence for dense clumps coincident with the shells. The expanding shells have an expansion speed of $\\sim$5 km s$^{-1}$. However, it is not clear whether HII region is the driving mechanism of the shell expansion.   We have presented a study about compact cores M and N located at the shell of nebula. Compact core N is associative with the H$_{2}$O maser, the OH maser, the IRAS point source, and the UC HII region. The spectral velocity components of the H$_{2}$O maser and the OH maser indicate that the star formation region is strongly activate. However, the deep-going properties of the compact cores M and N are needed to further investigate.   We also used the GLIMPSE, and the MIPSGAL survey data to analyze the YSOs and the warm dust distribution around the bubble N68. We identified class I and II YSOs using the [5.8]-[8.0] versus [3.6]-[4.5] relation, and correlated their distribution relative to the PDR, which we assume to be associated with and surrounding an HII region. We find that the N68 appears to have a significant number of YSOs associated with their PDRs, implying that triggered star formation mechanisms acting on the boundary of the expanding HII region."
    },
    "1208/1208.6563_arXiv.txt": {
        "abstract": "In this paper, we compare dispersions of a scalar field in Euclidean quantum gravity with stochastic inflation. We use Einstein gravity and a minimally coupled scalar field with a quadratic potential. We restrict our attention to small mass and small field cases.  In the Euclidean approach, we introduce the ground state wave function which is approximated by instantons. We used a numerical technique to find instantons that satisfy classicality. In the stochastic approach, we introduce the probability distribution of Hubble patches that can be approximated by locally homogeneous universes down to a smoothing scale. We assume that the ground state wave function should correspond to the stationary state of the probability distribution of the stochastic universe.  By comparing the dispersion of both approaches, we conclude three main results. (1) For a statistical distribution with a certain value, we can find a corresponding instanton in the Euclidean side, and it should be a complex-valued instanton. (2) The size of the universe of the Euclidean approach corresponds to the smoothing scale of the stochastic side; the universe is homogeneous up to the Euclidean instanton. (3) In addition, as the mass increases up to a critical value, both approaches break at the same time. Hence, generation of classical inhomogeneity in the stochastic approach and the instability of classicality in the Euclidean approach are related. ",
        "introduction": "In the semi-classical approach of gravity, we quantize fields in a classical metric background. Although the semi-classical approach cannot be the fundamental theory, it could give a useful guideline on how to approach to the problems of quantum gravity. One of the examples is Hawking radiation in black hole physics \\cite{Hawking:1974sw}. To study the nature of Hawking radiation, there are two useful approaches. One way is to treat quantum fluctuations of a field on the classical curved background \\cite{Birrell:1982ix}. The other way is to use the Euclidean path integral and obtain some thermodynamic quantities by the Wick rotation \\cite{Gibbons:1994cg}.  In the same spirit, we can tackle cosmological problems with two different approaches. As an application of the former (quantum field theory in a curved background), we can illustrate \\textit{stochastic inflation} \\cite{Vilenkin:1983}\\cite{Linde:1993xx}\\cite{Tolley:2008na}. Here, we can divide the Fourier modes into those shorter than the Hubble radius and longer than the Hubble radius. As long as the potential is broad and the field value is small, we can approximate that the longer wavelengths behave like a locally homogeneous and classical field while the shorter wavelengths act as a Gaussian random noise to the longer wavelengths. Therefore, the locally homogeneous scalar field will behave like a Brownian particle. This randomly walking field can be described by the \\textit{Langevin equation}. In the entire (inhomogeneous) universe, there are many Hubble patches that behave like this, and one may define the probability distribution of the fields for each of the Hubble patches. The master equation of this probability distribution is the \\textit{Fokker-Planck equation}.  As an application of the latter (Euclidean quantum gravity), it is useful to study the path integral. The ground state solution of the Schrodinger equation with gravity, so-called the \\textit{Wheeler-DeWitt equation}, is described by the Euclidean path integral \\cite{Hartle:1983ai}. It is not easy to calculate the whole path integral; however, if we only restrict to the on-shell solutions (instantons), then we can approximately obtain the wave function. In many contexts, instantons are useful to study non-perturbative phenomena of the universe. For example, the Coleman-DeLuccia instanton \\cite{Coleman:1980aw} is useful to describe the inhomogeneous vacuum decaying process, while the Hawking-Moss instanton \\cite{Hawking:1981fz} is used to describe the homogeneous tunneling process.  One interesting observation is that the stationary solution of the Fokker-Planck equation corresponds to the Hawking-Moss instanton \\cite{Linde:1993xx}. The question is that \\textit{is this correspondence an accident or not?} If there is no gravity, we can find many examples that a solution of the Fokker-Planck equation is indeed a solution of the Schrodinger equation \\cite{LeB}. However, it is not trivial if we include gravitation. Can we still extend this relation in the presence of gravity?  In this paper, keeping in mind this problem, we suggest three questions: \\begin{itemize} \\item If the stationary solution of the Fokker-Planck equation and the ground state of the superposition of the instantons are (approximately) the same, what is the Euclidean instanton that corresponds to a certain state for the stationary solution? \\item We already know that if the field is almost static, two approaches give the same results. Then, what will happen if the field begins to move? Even for this limit, are these two approaches the same? \\item If we increase the curvature of the potential (increase the mass around the local minimum), the stochastic description will breakdown. What is the corresponding phenomena in the Euclidean side? \\end{itemize} To investigate these issues, in Section~\\ref{sec:sto} and Section~\\ref{sec:euc}, we discuss the basics of the stochastic approach and the Euclidean approach. In Section~\\ref{sec:the}, we compare two approaches and discuss the first observations of both approaches. In Section~\\ref{sec:com}, we compare the details of two approaches, and we will answer the previous questions. Finally, in Section~\\ref{sec:con}, we summarize our discussions. ",
        "conclusions": ""
    },
    "1208/1208.3170_arXiv.txt": {
        "abstract": "{Embedded planets disturb the density structure of the ambient disk, and gravitational back-reaction possibly will induce a change in the planet's orbital elements. Low-mass planets only have a weak impact on the disk, so their wake's torque can be treated in linear theory. Larger planets will begin to open up a gap in the disk through nonlinear interaction. Accurate determination of the forces acting on the planet requires careful numerical analysis. Recently, the validity of the often used fast orbital advection algorithm (FARGO) has been put into question, and special numerical resolution and stability requirements have been suggested.  % } {We study the process of planet-disk interaction for low-mass planets of a few Earth masses, and reanalyze the numerical requirements to obtain converged and stable results. One focus lies on the applicability of the FARGO-algorithm. Additionally, we study the difference of two and three-dimensional simulations, compare global with local setups, as well as isothermal and adiabatic conditions. } {We study the influence of the planet on the disk through two- and three-dimensional hydrodynamical simulations. To strengthen our conclusions we perform a detailed numerical comparison where several upwind and Riemann-solver based codes are used with and without the FARGO-algorithm. } {With respect to the wake structure and the torque density acting on the planet, we demonstrate that the FARGO-algorithm yields correct a correct and stable evolution for the planet-disk problem, and that at a fraction of the regular cpu-time. We find that the resolution requirements for achieving convergent results in unshocked regions are rather modest and depend on the pressure scale height $H$ of the disk. By comparing the torque densities of two- and three-dimensional simulations we show that a suitable vertical averaging procedure for the force gives an excellent agreement between the two. We show that isothermal and adiabatic runs can differ considerably, even for adiabatic indices very close to unity. } {} ",
        "introduction": "\\label{sec:introduction}  Very young planets that are still embedded in the protoplanetary disk will disturb the ambient density by their gravity. This will lead to gravitational torques that can alter the orbital elements of the planet. For massive enough planets, the wake becomes nonlinear, and gap formation sets in. In numerical simulations of embedded planets, different length scales have to be resolved, in particular when studying low-mass planets. On the one hand, the global structure has to be resolved to be able to obtain the correct structure of the wakes, i.e. the spiral arms generated by the planet, which requires a sufficiently large radial domain. The libration of co-orbital material on horseshoe streamlines requires a full azimuthal extent of 2 $\\pi$ radians to be properly captured. On the other hand, the direct vicinity of the planet has to be resolved to study detail effects, such as horseshoe drag or accretion onto the planet.  To ease computational requirements, often planet-disk simulations are performed in the two-dimensional (2D) thin disk approximation, because a full three-dimensional (3D) treatment with high resolution is still very time-consuming. However, even under this reduced dimensionality, the problem is still computationally very demanding. The main reason is the strongly varying timestep size caused by the differentially rotating disk. Because the disk is highly supersonic with (azimuthal) Mach numbers of about 10 to 50, the angular velocity at the inner disk will limit the timestep of the whole simulation, even though the planet or other regions of interest are located much farther out. Changing to a rotating coordinate system will not help too much owing to the strong differential shear. To solve this particular problem and speed up the computation, \\citet{2000A&AS..141..165M} has developed a fast orbital advection algorithm (FARGO). This method consists of an analytic, exact shift in the hydrodynamical quantities by approximately the average azimuthal velocity. The transport step utilizes only the residual velocity, which is close to the local sound speed. Depending on the grid layout and the chosen radial range, a very large speed-up can be achieved, while at the same time the intrinsic numerical diffusion of the scheme is highly reduced \\citep{2000A&AS..141..165M,2000ASPC..219...75M}.  The original version of the algorithm has been implemented into the public code {\\tt FARGO}, which is very often used in planet-disk and related simulations. The accuracy of the FARGO-algorithm has been demonstrated in a detailed planet-disk comparison project utilizing embedded Neptune and Jupiter mass planets \\citep{2006MNRAS.370..529D}. There, it has been shown that it leads to identical density profiles near the planet and total torques acting on the planet. Meanwhile, similar orbital advection algorithms have been implemented into a variety of different codes in two and three spatial dimensions, e.g. {\\tt NIRVANA} \\citep{1997CoPhC.101...54Z,2009A&A...506..971K}, {\\tt ATHENA} \\citep{2008JCoPh.227.4123G,2010ApJS..189..142S}, and {\\tt PLUTO} \\citep{2007ApJS..170..228M, Pluto2012}. Despite these widespread applications, it has been claimed recently that usage of the FARGO-algorithm (here in  connection with {\\tt ATHENA}) may lead to an unsteady behavior of the flow near the planet, which even affects the wake structure of the flow farther away from the planet \\citep{2011ApJ...741...56D}.  In the same paper, \\citet{2011ApJ...741...56D} note that a very high numerical grid resolution is required to obtain a resolved flow near the planet. In particular, they analyze the smooth, unshocked wake structure close to the planet and infer that a minimum spatial resolution of about 256 gridcells per scale height $H$, of the disk is needed to obtain good agreement with linear studies. New simulations with a moving mesh technique also seem to indicate the necessity of very high resolutions \\citep{2012ApJ...755....7D}.  Because a robust, fast, and reliable solution technique is mandatory in these type of simulations, we decided to address the planet-disk problem for a well defined standard setup, which is very close to the one used in \\citet{2011ApJ...741...56D}. To answer the question of the validity of the FARGO-algorithm and estimate the resolution requirements, we applied several different codes to an identical problem. These range from classical second-order upwind schemes (e.g. {\\tt RH2D}, {\\tt FARGO}) to modern Riemann-solvers such as {\\tt PLUTO}. The characteristics of these codes are specified in Appendix \\ref{seca:codes}.  Another critical issue in planet-disk simulations is the selection of a realistic treatment of the gravitational force between the disk and the planet. Because the planet is typically treated as a pointmass and located within the numerical grid, regularization of the potential is required. In addition, physical smoothing is required to account for the otherwise neglected vertical thickness of the disk. The magnitude of this smoothing is highly relevant, since it influences the torques acting on the planet \\citep{2002A&A...387..605M}, and the smoothing parameter has even entered analytical torque formulas \\citep{2009ApJ...703..857M,2010MNRAS.401.1950P}. Because the 2D equations are obtained by a vertical averaging procedure, the force should be calculated by a suitable vertical integration as well. This approach has been undertaken recently by \\citet{2012A&A...541A.123M}, who show that the smoothing length is indeed determined by the vertical thickness of the disk, and is roughly on the order of $0.7\\,H$. They show in addition that the change of the disk thickness induced by the presence of the planet has to be taken into account. Because in recent simulations very short smoothing lengths have been used in 2D simulations \\citep{2011ApJ...741...56D,2012ApJ...755....7D}, we compare our 2D results on the standard problem to an equivalent 3D setup and infer the required right amount of smoothing.  Finally, we performed additional simulations for different equations of state. The first set of simulations deals with the often used locally isothermal setup, while in comparison simulations we explore the outcome of adiabatic runs. This is important because some codes may not allow for treating an isothermal equation of state. Here, we use different values for the ratio of specific heat $\\gamma$. In particular, a value of $\\gamma$ very close to unity has often been quoted as closely resembling the isothermal case. We show that this statement can depend on the physical problem. In particular, in flows where the conservation of the entropy along streamlines is relevant, there can be strong differences between an isothermal and an adiabatic flow, regardless of the value chosen for $\\gamma$. For the planet-disk problem this has already been shown by \\citet{2008A&A...478..245P}.  In the following section \\ref{sec:model} we describe the physical and numerical setup of our standard model, and present the numerical results in Sect.~\\ref{sec:std-results}. The validity of the FARGO-algorithm is checked in Sect.~\\ref{sec:numerics}. Alternative setups (nearly local, 2D versus 3D, adiabatic) are discussed in Sect.~\\ref{sec:alter}. The transition of the wake into a shock front is discussed in Sect.~\\ref{sec:shock}, and in the last section, we summarize our results. ",
        "conclusions": "\\label{sec:summary} Through a series of 2D and 3D simulations using different computational methods and codes we have explored in detail the numerical requirements for studies of the planet-disk problem. In our analysis we focus on the torque density acting on the planet and the structure of the wake generated by the planet.  With respect to the applicability of the fast orbital advection algorithm, FARGO, we have shown that it leads to consistent numerical results that agree extremely well with non-FARGO studies. The achievable gain in speed can be significant. For the setup used here we found a speed-up of more than a factor of 10. The method works well in the presence of embedded planets, does not show any signs of unsteady behavior, and can be applied in two or three spatial dimensions. As it is applicable in conjunction with magnetic fields as well, new possibilities with respect to numerical studies of turbulent accretion disks open up \\citep{Pluto2012}.  Concerning the treatment of the gravitational potential of embedded planets, we extend previous studies \\citep{2002A&A...387..605M,2012A&A...541A.123M} to very low mass planets in extremely thin disks. We confirm that, for physical reasons, in 2D simulations the planetary potential has to be smoothed with about $\\epsilon = 0.6\\,H $ -- $0.7\\,H$. Models where the gravitational force is obtained directly through a vertical integration yield always reasonable agreement with full 3D simulations. The usage of very small smoothing lengths below $\\epsilon = 0.6\\,H$ in 2D simulations is not recommended, because then the forces in the vicinity of the planet are strongly overestimated, which results in an unphysical enhancement of the torque and too strong wakes.  Through a careful resolution study, we show that the smooth wake structure at distances smaller than about $2\\,H$ of the planet can be resolved well, and consistently, already with very low resolution of 8 to 16 cells per scale height. The results are clearly converged for 32 grid-cells per $H$. For larger distances from the planet, the spiral wake turns into a shock wave and much higher resolution may be required. We found, that around a resolution of about 100 grid-cells per $H$ convergence can be achieved. Because this high resolution is only required near the spiral shocks, and the flow is relatively smooth outside, numerical methods that adaptively refine this crucial region may be the method of choice in the future.  For adiabatic flows we confirm earlier findings \\citep{2008A&A...478..245P} that the unsaturated horseshoe drag shows a strong deviation from the isothermal case. Using the appropriate scaling the adiabatic corotation torques are independent of $\\gamma$ and do not converge to the isothermal case, even in the limit $\\gamma \\rightarrow 1$. Hence, the procedure of modelling the isothermal case with simulations of $\\gamma$ close to unity, has to be treated with care. In the final saturated case, where all the corotation effects have been wiped out, isothermal and adiabatic results agree perfectly, once the correction to the sound speed has been applied.  In Appendix \\ref{subsec:std-resolution} we have shown that we do not find an additional timestep criterion due to the planetary potential and we also have not noticed any unstable evolution in the case of using the orbital advection. The question why in the simulations using the {\\tt ATHENA}-code instabilities occur \\citep{2011ApJ...741...56D} may be connected to the treatment of orbital advection in that code \\citep{2010ApJS..189..142S} which is apparently different from the implementation in the {\\tt FARGO}-code. One should also notice that in such simulations the conservative treatment of Coriolis forces is mandatory to properly conserve angular momentum \\citep{1998A&A...338L..37K}.  We have demonstrated that the planet-disk interaction problem may be regarded as a very good test to validate an implementation of orbital advection, because it admits a nearly analytic solution to which a code output can be compared. This is not the case for simulations of turbulent disks, where no such known solutions exist. We hope, that the presented results and comparison simulations may serve as a useful reference for other researcher in this field."
    },
    "1208/1208.5560_arXiv.txt": {
        "abstract": "We present the Suzaku results of the supernova remnant (SNR) G346.6$-$0.2. The X-ray emission has a center-filled morphology with the size of 6$'$$\\times$8$'$ within the radio shell. Neither an ionization equilibrium nor non-equilibrium (ionizing) plasma can reproduce the spectra remaining shoulder-like residuals in the 2--4 keV band. These structures are possibly due to recombination of free electrons to the K-shell of He-like Si and S. The X-ray spectra are well fitted with a plasma model in a recombination dominant phase. We propose that the plasma was in nearly full ionized state at high temperature of $\\sim$5 keV, then the plasma changed to a recombining phase due to selective cooling of electrons to lower temperature of $\\sim$0.3 keV. G346.6$-$0.2 would be in an epoch of the recombining phase. ",
        "introduction": "G346.6$-$0.2 is a supernova remnant (SNR) discovered in the radio band \\citep{Clark1975}. The radio image shows a shell structure with the size of $\\sim$8$'$ \\citep{Clark1975,Dubner1993, Whiteoak1996}. Flux densities at 408 MHz, 843 MHz, 1.47 GHz, and 5 GHz were measured to be 14.9 Jy, 8.7 Jy, 8.1 Jy, and 4.3 Jy, respectively (\\cite{Green2009} and references therein), then the spectral index was estimated to be 0.5 \\citep{Clark1975,Green2009}. The SNR would be interacted with a molecular cloud because an OH maser was found \\citep{Green1997}.  The Galactic plane survey project with ASCA discovered a faint X-ray emission from G346.6$-$0.2 for the first time \\citep{Yamauchi2008}. The ASCA GIS image showed the diffuse X-ray morphology in the radio shell. The X-ray spectrum was represented by either a thermal plasma model with a temperature of $\\sim$1.6 keV or a power-law model with a photon index of $\\sim$3.7. The absorption was as large as $N_{\\rm H}=$(2--3)$\\times10^{22}$ cm$^{-2}$ \\citep{Yamauchi2008}, which suggested that the SNR is located at a long distance, possibly in the Galactic inner disk or further. No further quantitative constraint was available with the ASCA data due to the limited photon statistics. G346.6$-$0.2 was then observed with Suzaku. \\citet{Sezer2011} found strong emission lines of Si and S, and fitted the X-ray spectrum with a model of a power-law (photon index $\\sim$0.6) plus a non-equilibrium ionization (NEI) (ionizing) plasma. They predicted that the strong Si and S lines are due to an ejecta-dominated plasma which originated from a Type Ia supernova explosion, and the power-law component is regarded as synchrotron emission.  Recently, strong radiative recombination continua (RRCs) have been discovered in the X-ray spectra of five mixed-morphology (MM) SNRs \\citep{Yamaguchi2009,Ozawa2009,Ohnishi2011, Sawada2012, Uchida2012}. The RRC originates from radiative transitions of free electrons to the K-shell of ions, a sign of a recombination dominant plasma (RP). In the residuals of the NEI model fit for G346.6$-$0.2 in \\citet{Sezer2011}, we see a similar structure to the RRC.  G346.6$-$0.2 is located on the Galactic ridge, where a strong X-ray emission, called the Galactic Ridge X-ray Emission (GRXE), is prevailing. However, the previous data reduction and analyses did not properly take account of the GRXE as a major background for the faint and diffuse source. We, therefore, revisited the Suzaku data and performed the data reduction, spectral construction, and spectral analysis paying particular concerns on the subtraction of the GRXE. We then discovered evidence for the recombining plasma from this MM SNR for the first time. ",
        "conclusions": "\\subsection{G346.6$-$0.2}  For the background estimation, we took account of the difference of vignetting effects between the source and background regions. Furthermore, we tried two different backgrounds of BGD-a and BGD-b (spectra-a and spectra-b). Although the fluxes of these backgrounds are different, the spectral shapes are very similar. As the results, these two spectra, spectra-a and spectra-b, gave essentially the same best-fit parameters except for the absolute luminosity. Since BGD-a and G346.6$-$0.2 are likely in a local soft X-ray excess around  G346.6$-$0.2, we discuss based on the results of spectra-a, the BGD-a subtracted spectra.  The spectrum was nicely described by the RP in a transition epoch of its recombining phase ({\\tt neij} model). The best-fit electron temperature was about 0.3~keV. These conclusions are inconsistent with those of \\citet{Yamauchi2008} and \\citet{Sezer2011}. In the thermal plasma fitting, their results were in CIE or NEI with the electron temperatures of $\\sim$1.6 keV and 0.97--1.3 keV in \\citet{Yamauchi2008} and \\citet{Sezer2011}, respectively. These apparent inconsistencies are mainly due to the background subtraction. \\citet{Sezer2011} used the background from very small region (near the corner of the FOV), while \\citet{Yamauchi2008} used the annulus region around the source. No vignetting effect was corrected in \\citet{Yamauchi2008} and \\citet{Sezer2011}. Therefore, the background-subtracted spectra provide less statistics and should contain some fractions of the GRXE emission, particularly in the hard X-ray band. In fact, they reported that the X-ray spectrum has excess flux above $\\sim$5 keV. This may artificially predict higher electron temperatures and/or power-law component.  Assuming the mean density of 1H cm$^{-3}$, the derived $N_{\\rm H}$ value of (2.3$\\pm0.1$)$\\times$10$^{22}$ cm$^{-2}$ corresponds to the distance of 7--8 kpc, which is well consistent with the distance of 8.2 kpc estimated from $\\Sigma$-$D$ relation \\citep{Case1998}. If we assume the distance of 8 kpc, the luminosity was calculated to be 3.5$\\times10^{35}$ erg s$^{-1}$ in the 0.5--10 keV energy band. The abundances are sub-solar to solar, which is consistent with those of \\citet{Sezer2011}. They predicted that the remnant may originate from a Type Ia supernova (SN) explosion, and  the solar/sub-solar abundances of heavy elements are explained by the scenario that the SNR is young, and the reverse shock does not reach yet to the interior of the ejecta. Our best-fit $n_{\\rm e} t$ is (4.8$^{+0.1}_{-0.4}$)$\\times10^{11}$ cm$^{-3}$ s. Then, assuming the mean density of 1H cm$^{-3}$, the age of G346.6$-$0.2 is estimated to be 1.4$\\times10^4$--1.6$\\times10^4$ yr. This means that G346.6$-$0.2 is not a young SNR, and hence no evidence for the young Type Ia SN scenario is obtained. The solar/sub-solar abundances plasma would be mainly due to interstellar matter.  G346.6$-$0.2 is the sixth sample of the SNRs with the RP, after IC443 \\citep{Yamaguchi2009}, W49B \\citep{Ozawa2009}, G359.1$-$0.5 \\citep{Ohnishi2011}, W28 \\citep{Sawada2012}, and W44 \\citep{Uchida2012}. Among these SNRs with the RP, the plasma structures in W28 and W44 were studied in detail. The results were that the RP is in a recombining phase of $\\sim 10^{11}$($n_{\\rm e}/1{\\rm H~cm}^{-3}$)$^{-1}$~s after the production of the initial RP.  We found that G346.6$-$0.2 is explained with a similar scenario as these SNRs. We propose that the other SNRs, IC443, W49B, and G359.1$-$0.5, can also be explained by the same scenario, a plasma in an epoch of recombining process.  All the previous RP-detected SNRs share some common characteristics: (1) they are classified to MM SNRs \\citep{Rho1998}, (2) an OH maser is detected, suggesting the interaction with molecular clouds, and (3) TeV/GeV $\\gamma$-ray emission is detected. For possible origins of the RP based on these common features, one can refer the discussions in \\citet{Yamaguchi2009}, \\citet{Sawada2012}, and \\citet{Uchida2012}. The center-filled thermal X-ray emission within the radio shell suggests that G346.6$-$0.2 is a MM SNR. An OH maser has been found \\citep{Green1997}, but no TeV/GeV $\\gamma$-ray emission has been found from G346.6$-$0.2. Therefore, future TeV/GeV $\\gamma$-ray search from this SNR is encouraged.  \\subsection{Local Emission }  The excess emission around G346.6$-$0.2 was shown by the optically thin thermal plasma with the temperature of 0.79 keV, 0.14 solar abundance, and $N_{\\rm H}$ value of 1.2$\\times10^{22}$ cm$^{-2}$. The ASCA Galactic plane survey revealed that the GRXE spectra observed in various regions are well represented by an optically thin thermal plasma model with two temperatures of $<$1 keV and $\\sim$7 keV \\citep{Kaneda1997a,Kaneda1997b}. The two-temperature structure was confirmed by Chandra \\citep{Ebisawa2005} and Suzaku \\citep{Ryu2009}. The spectral parameters of the excess emission in the G346.6$-$0.2 field are similar to those of the soft component of the GRXE. The intensity distribution of the hard component along the Galactic plane is symmetric with respect to the Galactic center, while that of the soft component is more asymmetric: a local peak and a local minimum were found at $l\\sim$347$^{\\circ}$ and $l\\sim$355$^{\\circ}$, respectively \\citep{Kaneda1997a}; the longitudinal distribution of the GRXE soft component shows that the intensity near the BGD-a region is about $\\sim$1.3 times higher than that of the BGD-b region \\citep{Kaneda1997a}. Our present result of BGD-a is $\\sim$2 times larger than BGD-b  in the soft X-ray band, and hence the excess would not be due to the fluctuation of the GRXE, but would be a local plasma. The $N_{\\rm H}$ value of the local plasma is smaller than that of G346.6$-$0.2 in the GRXE, hence this emission is located at the near side of G346.6$-$0.2 and the GRXE. The local plasma is near to the direction of the non-thermal SNR RX J1713.7$-$3946 and the $N_{\\rm H}$ value is similar to that of RX J1713.7$-$3946 \\citep{Koyama1997}. Therefore, the local plasma is located near RX J1713.7$-$3946 at the distance of 1--2 kpc. Whether this thermal plasma is physically associated with RX J1713.7$-$3946 or not is unclear. To clarify this, we encourage further observations around RX J1713.7$-$3946.  \\bigskip  We would like to express our thanks to all of the Suzaku team. We thank Dr. M. Sawada for his useful comments. This work was supported by the Japan Society for the Promotion of Science (JSPS); the Grant-in-Aid for Scientific Research (C) 21540234 (SY), 24540232 (SY), and 24540229 (KK), Young Scientists (B) 24740123 (MN), Challenging Exploratory Research program  20654019 (KK), and Specially Promoted Research 23000004 (KK)."
    },
    "1208/1208.3895_arXiv.txt": {
        "abstract": "{Since the discovery of the first exoplanet in 1995 around a solar-type star, the interest in exoplanetary systems has kept increasing. Studying exoplanet host stars is of the utmost importance to establish the link between the presence of exoplanets around various types of stars and to understand the respective evolution of stars and exoplanets.} {Using the limb-darkened diameter (LDD) obtained from interferometric data, we determine the fundamental parameters of four exoplanet host stars. We are particularly interested in the F4 main-sequence star, $\\theta$~Cyg, for which Kepler has recently revealed solar-like oscillations that are unexpected for this type of star. Furthermore, recent photometric and spectroscopic measurements with SOPHIE and ELODIE (OHP) show evidence of a quasi-periodic radial velocity of $\\sim$ 150 days. Models of this periodic change in radial velocity predict either a complex planetary system orbiting the star, or a new and unidentified stellar pulsation mode.} {We performed interferometric observations of $\\theta$~Cyg, 14~Andromedae, $\\upsilon$ Andromedae and 42~Draconis for two years with VEGA/CHARA (Mount Wilson, California) in several three-telescope configurations. We measured accurate limb darkened diameters and derived their radius, mass and temperature using empirical laws.} {We obtain new accurate fundamental parameters for stars 14~And, $\\upsilon$ And and 42~Dra. We also obtained limb darkened diameters with a minimum precision of $\\sim1,3\\%$, leading to minimum planet masses of $M \\sin i = 5.33\\pm0.57$, $0.62\\pm0.09$ and $3.79\\pm0.29$ $M_{\\rm Jup}$ for 14~And~b, $\\upsilon$~And~b and 42~Dra~b, respectively. The interferometric measurements of $\\theta$~Cyg show a significant diameter variability that remains unexplained up to now. We propose that the presence of these discrepancies in the interferometric data is caused by either an intrinsic variation of the star or an unknown close companion orbiting around it.}  {} ",
        "introduction": "\\label{sect:Intro}  Many techniques have been developed during the past decade to enable the discovery of exoplanets. The radial velocity method, based on the reflex motion of the host star, is one of the most successful of these and has to date enabled the discovery of 535 planetary systems\\footnote{As of December 23, 2011 \\citep{encyclopedia}}. Most of these planets were found orbiting slowly rotating stars, late-type stars, or A giants. While A and F main sequence stars were usually avoided because of their high $v\\sin i$, a survey of A and F main sequence stars was nonetheless recently undertaken using a specialized analysis method to look for planets around these stars, and planets were indeed found around a few F stars \\citep{Lagrange2009}. Unfortunately, the possible planet configurations fitting the radial velocity (RV) data were found to be dynamically unstable. To resolve this problem it is important to better understand the link between the presence and mass of exoplanets, the host star parameters, and the separation of the planet and host star.  Interferometric data are now able to bring additional information to bear on stellar variability and its contribution to noise in the radial velocity measurements, and can help to directly determine many of the fundamental parameters of the host stars with an accuracy of about $5 \\%$ see for example \\cite[see for example][]{baines,55cnc}. This is not only very important for deriving accurate radii for transiting planets, but also for RV planets. Understanding the link between the presence and nature of exoplanets and the fundamental parameters of the star requires sampling a large number of targets. We have started a survey with VEGA (Visible spEctroGraph and polArimeter) \\citep{Vega1}, a visible spectro-interferometer located on the CHARA (Center for High Angular Resolution Astronomy) array at Mount Wilson, California \\citep{chara}, to measure all currently accessible exoplanets stars, i.e. almost 40 targets. To build this sample, we first selected the exoplanet host stars listed in Schneider's catalog \\citep{encyclopedia}. Those stars have to be observable by VEGA, therefore we sorted out those that had a magnitude smaller than $6.5$ in the V- and in the K band, and a declination higher than $-30^{\\circ}$. Knowing the error on the squared visibility allowed by VEGA at medium resolution ($\\simeq 2\\%$), we can estimate the maximum and the minimum diameters for which we can obtain an accuracy of $\\simeq 2\\%$ taking into account the maximum and minimum baselines. We consider this accuracy as the minimum allowed to obtain sufficiently good informations on fundamental parameters of the stars and planets. Diameters included between $0.3$ and $3$ milliseconds of arc (mas) are sufficiently resolved to achieve this accuracy. We finally found 40 stars whose planets were discovered with the transit or RV techniques.  Interferometry is complementary to the transit method or RV measurements in determining exoplanet parameters. For instance, the transit method allows determining the exoplanet radius, while the RV method is used to detect the minimum mass. The main goal of these observations is to directly constrain these parameters, and to study the impact of stellar noise sources (e.g., spots, limb darkening) applied to these observing methods. In the long term, the results will be compared to a catalog   of limb darkening laws from 3D hydro-\u00addynamical modeling and radiative transfer. Thus, we will be able to create a catalog of measured angular diameters, and derive revised surface brightness relationships.  From October to December 2011, we obtained data on three stars of our sample\\,: 14~And, $\\upsilon$ And and 42~Dra, while a fourth star $\\theta$~Cyg was observed over a longer period, from June 2010 to November 2011. We found that while the first three stars yield stable and repeatable results, there are  discrepancies in the results of $\\theta$~Cyg, forcing us to study this system more carefully. New and unexplained RV variations recorded with SOPHIE and ELODIE at the Observatoire de Haute-Provence \\citep{13cygDesort} provided a first clue that this star hosts either a complex planetary system, undergoes hitherto unknown variations, or has a hidden companion.  After a short introduction to the basics of interferometry, we describe in Section~\\ref{sect:exoplanets} the observations made of 14~And, $\\upsilon$ And and 42~Dra during the year 2011 and derive the star and planet fundamental parameters.  We then compare these values to those found in the literature. In Section~\\ref{sect:observations}, we present the observations of $\\theta$~Cyg made during the last two years. We discuss the fundamental parameters we derived for this target in Section~\\ref{sect:parameters}, and compare them with the previously known parameters for this star (see Table~\\ref{tab:table1}). We then discuss the variation of the angular diameter of $\\theta$~Cyg in Section~\\ref{sect:discussion} and some possible explanations of this variability. ",
        "conclusions": "We obtained new and accurate visibility measurements of 14~And, $\\upsilon$ And and 42~Dra using visible band interferometric observations. From these we derived accurate values of the LD diameter and of fundamental parameters that are fully consistent with those derived with other techniques and bring some improvements in precision. The error bars and $\\chi^2_{\\rm reduced}$ for these three stars are in general much smaller than those obtained on our fourth target\\,: $\\theta$~Cygni. We analyzed the scatter of measurements of $\\theta$~Cyg, taking into account that instrumental or data processing bias are well understood thanks to the good results obtained on the three other stars. It appears that a solution with an unknown companion close to the star helps in reducing the residuals in the model fitting. The limited accuracy in our determination prevents us from being conclusive about the presence of a new close companion around $\\theta$~Cyg, and do not allow us to tell which type of star it could be, because it is not necessarily bound. However, this result encourages organizing new observations in the visible and IR wavelengths, focused on confirming or denying this hypothesis. Closure phase signal is a good way to detect and characterize faint companions around bright stars. We performed simulations of expected closure phase signal for $\\theta$~Cyg with a companion contributing $10\\%$ of the flux that is located at $25$ mas. As explained before, VEGA is unfortunately not able to measure accurate closure phase signals. CLIMB in K band and MIRC \\citep{mirc} in H band are well-adapted for closure phase tests with the largest CHARA triangle (E1W1S1). Expected signals are presented in Fig.~\\ref{fig:cloture}. Therefore, a more adequate observing strategy and dedicated observations will be prepared with the combination of the different CHARA beam combiners.  \\begin{figure} \\centering \\includegraphics[width=8cm,height=4cm]{./clotureclimb.eps}\\\\ \\includegraphics[width=8cm,height=4cm]{./cloturemirc.eps} \\caption{Closure phase signals expected for $\\theta$~Cyg with a companion located at $25$ mas that contributes to $10\\%$ of the flux. UP\\,: CLIMB in K band, BOTTOM\\,: MIRC in H band. We used the largest triangle of telescopes on CHARA\\,: E1W1S1.} \\label{fig:cloture} \\end{figure}"
    },
    "1208/1208.4198_arXiv.txt": {
        "abstract": "{ We have computed the Galactic Habitable Zones (GHZs) of the Andromeda galaxy (M31) based on the probability of terrestrial planet formation, which depends on the metallicity ($Z$) of the interstellar medium, and the number of stars formed per unit surface area. The GHZ was therefore obtained from a chemical evolution model (CEM) built to reproduce a metallicity gradient in the galactic disk, [O/H]($r$) $ = -0.015 \\pm 0.003$ dex kpc$^{-1} \\times r$(kpc) $ + 0.44 \\pm 0.04$ dex. This gradient is the most probable when intrinsic scatter is present in the observational data. The chemical evolution model predicted a higher star formation history (SFH) in both the halo and disk components of M31 and a less efficient inside-out galactic formation, compared to those of the Milky Way. If we assumed that Earth-like planets form with a probability law that follows the $Z$ distribution shown by stars with detected planets and the SFH predicted by the CEM, the most probable GHZ per $pc^2$ is located between  3 and 7 kpc for planets with ages between 6 and 7 Gy, approximately. But the highest number of stars with habitable planets is in a ring located between 12 and 14 kpc with mean age of $\\sim$ 7 Gy. 11 \\% and 6.5 \\% of the all formed stars in M31 may have planets capable of hosting basic and complex life, respectively. } \\resumen { Calculamos la Zona de Habitabilidad Gal\\'actica (ZHG) de M31 bas\\'andonos en la probabilidad de formaci\\'on de planetas terrestres dependiente de la metalicidad ($Z$) del medio interestelar. La ZHG fue determinada a partir de un modelo de evoluci\\'on qu\\'{\\i}mica construido para reproducir un gradiente de metalicidad  en el disco gal\\'actico: [O/H]($r$)$ =-0.015$ dex kpc$^{-1} \\times r$(kpc)$ + 0.44$ dex. Suponiendo que los planetas tipo Tierra se forman bajo una ley de probabilidad que sigue la distribuci\\'on de $Z$ mostrada por las estrellas con planetas, entonces la ZHG m\\'as probable se localiza entre 3 y 7 kpc para planetas con edades entre 6 y 7 Ga, aproximadamente, aunque el mayor n\\'umero de estrellas con planetas habitables se encuentra en un anillo localizado entre 12 y 14 kpc y edad promedio de $\\sim$ 7 Ga. El 11 \\% y 6.5 \\% de todas las estrellas formadas en M31 podr\\'{\\i}an albergar vida b\\'asica y compleja, respectivamente. } \\addkeyword{Galactic habitability} \\addkeyword{Chemical evolution} \\addkeyword{Abundance gradients} \\addkeyword{Andromeda galaxy} \\addkeyword{M31}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  The Galactic Habitable Zone (GHZ) is defined as the region with sufficient abundance of chemical elements to form planetary systems in which Earth-like planets could be found and might be capable of sustaining life (Gonzalez et al. 2001, Lineweaver 2001). An Earth-like planet is a rocky planet characterized in general terms by the presence of water and an atmosphere (Segura \\& Kaltenegger 2009).  GHZ research has focused mainly on our galaxy, the Milky Way (MW) (Gonzalez et al. 2001, Lineweaver et al. 2004, Prantzos 2008, Gowanlock et al. 2011). Gonzalez et al. were the first to propose the concept of a GHZ, which is a ring located in the thin disk that migrates outwards with time, due to their most important assumption is that terrestrial planet form with metallicities higher than 1/2 of the solar value ($Z_\\odot$). Lineweaver et al. (2004) later proposed that the Milky Way's GHZ is a ring located in the Galactic disk within a radius interval of 7 to 9 kpc from the center of the MW, and that the area of the ring increases with the age of the Galaxy, because they consider a $Z$ distribution, for forming terrestrial planets, that peaks at $\\sim 0.8 Z_\\odot$; the presence of a host star; and the absence of nearby supernovae (SN) harmful to life. On the other hand, Prantzos concluded that the current GHZ covers practically the entire MW disk, due to he assumes a $Z$ probability, to form Earth-like planets, almost equal for $Z > 0.1 Z_\\odot$. These studies confirmed that the Solar System is located within the GHZ since the Sun is found at 8 kpc from the center of the Galaxy. Nevertheless, a star with an Earth-like planet capable of sustaining life is more likely to be found in inner rings of the Galactic disk, between 2 and 4 kpc, owing to the high stellar surface density that is present in the inner disk of our galaxy (Prantzos 2008, Gowanlock et al. 2011).  The GHZ has recently been computed for two elliptical galaxies (Suthar \\& McKay 2012). Imposing only metallicity restrictions for planet formation, these authors found that both elliptical galaxies could sustain broad GHZs.  Here, we extend MW studies to the disk of the most massive galaxy in the Local Group: the Andromeda galaxy (M31). M31 is a type SAb spiral galaxy, whose visible mass is $\\sim$1.2 times larger than that of the MW, and M31 is at a distance of 783 $\\pm$ 30 kpc (Holland 1998).  The GHZ depends mainly on the abundance of chemical elements heavier than He (metallicity, $Z$), since $Z$ leads to planetary formation. Moreover,  the GHZ also depends on the occurrence of strong radiation events that can sterilize a planet. Melott \\& Thomas (2011) study many kinds of astrophysical radiations lethal to life, as electromagnetic radiation (e.g. X rays), high-energy protons or cosmic rays from stars (included the Sun), SN, and gamma-ray bursts (GRB). According to them, the SN and GRB could be the dominant cause of extinctions. Since most of the SN and GRB are originated during the last stages of massive stars, the rate of high-mass stars is useful to estimate the galactic zones where life on planets is annihilated by astrophysical events. In this paper, we excluded the bulge of Andromeda from the GHZ, despite its having a high enough abundance of chemical elements to form planets (Sarajedini \\& Jablonka 2005). The bulge, located between 0 and 3 kpc from de galactic center, might not provide a stable habitat for life, due to its high supernovae rate at early times, which could sterilize planets, and the proximity of stars that can destabilize the orbits of the planets (Jim\\'enez-Torres et al. 2011).  Throughout this paper the terms \"evolved life\" and \"complex life\" are synonymous. These terms refer to a type of life similar to that of the human beings, since they are able to develop advanced technologies.  In this study, we present a chemical evolution model (CEM) for the halo and disk components that predicts the temporal behavior of the space distribution of $Z$ and SN occurrence in M31, built on precise observational constraints (Sec. 2). Based on a set of biogenic, astrophysical, and geophysical restrictions, the CEM results lead to the determination of the GHZ (Sec. 3). We discuss the implication of the chemical evolution model and the GHZ condition on the location, size, and age of the Galactic Habitable Zone (Sec. 4). Finally, we present our conclusions of the present study (Sec. 5). ",
        "conclusions": "\\label{sec:thend}  Based on O/H values of {\\hii} regions, we obtained current O/H gradients with different empirical and theoretical calibrations. In the presence of intrinsic scatter, we computed the most probable gradient for M31 from theoretical calibrations based on the R$_{23}$ method and we found that [O/H]($r$) = $- 0.015 \\pm 0.003 $dex kpc$^{-1} \\times r(kpc) + 0.44 \\pm 0.04$ dex. The slope of the gradient is 0.03 dex kpc$^{-1}$ flatter, and the value at the center of the galaxy is 0.06 dex higher, than those values of the O/H gradient of the Milky Way disk.  The chemical evolution model built to reproduce our O/H gradient of the galactic disk matches the current radial distribution of the gas mass, $\\Sigma_{gas}(r)$, for the outer regions quite well, but it fails for the inner ones. On the other hand, the current radial distribution of the star formation rate, SFR($r$), fits for the inner parts, but not for the outer ones. Therefore, in order to improve the agreements, the model would require more complex galactic and star formation histories.  The model cannot reproduce the O/H gradient computed by empirical methods, which is 0.42 dex lower for the central value than the gradient obtained by theoretical methods, unless the other observational constraints are considerably modified.  Based on the chemical evolution model we obtained, for the first time, the Galactic Habitable Zone per $pc^2$ (GHZ) of M31, considering three space requirements (per surface unit):  i) sufficient metallicity for planet formation with a probability law that follows the $Z$ distribution shown by exoplanets, ii) high number of stars that may be potential home for life,  and iii) an average SN rate similar to that permitting the existence of life on the Earth; and two time requirements: the existence of basic life (like cyanobacteria), and the development of evolved life (like humans).  The GHZ of M31 with high probability is located between 3 and 7 kpc on planets with ages between 6 and 7 Gy, approximately. Assuming the area of each ring, the maximum number of stars, of all ages and harboring Earth-like planets, is in a ring located between 12 and 14 kpc with mean ages of $\\sim$ 6 Gy and $\\sim$ 7 Gy for planet capable of sustaining basic and complex life, respectively.  The GHZ of M31 with high and medium probability is located between 3 and 14 kpc (the inner half of the disk) on planets with ages between 3 and 9 Gy, approximately. The width at half maximum of number of stars, of any ages and harboring Earth-like planets suitable for basic life, is in an annular region located between 9 and 18 kpc, with mean age of $\\sim$ 5 Gy. That annular region corresponds to 27 \\% of the galactic disk area ($\\sim \\pi (30^2 - 3^2) kpc^2$). On the other hand, that width for complex life is located between 10 and 18 kpc, which corresponding to  25 \\% of the galactic disk, and the mean age of  that annular region is $\\sim$ 6.5 Gy. According our computation 11 \\% of the stars formed in M31 may have planets capable of hosting basic life, and 6.5 \\% complex life.  SN effects are important only in the inner half disk and during the first half of the galactic evolution, but $Z$ effects are more important during the second half of the galactic evolution.  In GHZ studies, the $Z$ restriction is crucial for finding the location of Earth-like planets, but the SN survival condition allows us to compute the location of those planets that survive SN events.  Based on the previous GHZs of the MW and the present GHZ of the M31, we are not able to compare the GHZs of those spiral galaxies, due to the GHZs were computed with different  constraints.  \\vskip 1cm  Part of this work was submitted by Sof\\'{\\i}a Meneses Goytia to the Master's Programme in Chemical Sciences at the Universidad Nacional Aut\\'onoma de M\\'exico."
    },
    "1208/1208.1297_arXiv.txt": {
        "abstract": "It has recently been proposed that radio emission from magnetars can be evaluated using a ``fundamental plane'' in parameter space between pulsar voltage gap and ratio of X-ray luminosity $\\Lx$ to rotational energy loss rate $\\Edot$. In particular, radio emission from magnetars will occur if $\\Lx/\\Edot<1$ and the voltage gap is large, and there is no radio emission if $\\Lx/\\Edot>1$. Here we clarify several issues regarding this fundamental plane, including demonstrating that the fundamental plane is not uniquely defined. We also show that, if magnetars and all other pulsars are different manifestations of a unified picture of neutron stars, then pulsar radio activity (inactivity) appears to be determined by the ratio $\\Lx/\\Edot\\lesssim1$ ($\\Lx/\\Edot\\gtrsim 1$), although observational bias and uncertainty in the ratio for some sources may still invalidate this conclusion. Finally, we comment on the use of other pulsar parameters that are constructed from the three observables: spin period $P$, period derivative $\\Pdot$, and $\\Lx$. ",
        "introduction": "\\label{sec:intro}  Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) form the magnetar class of neutron stars, i.e., neutron stars which possess superstrong magnetic fields ($B\\gtrsim 10^{14}\\mbox{ G}$) in most cases. Their strong fields likely power the activity seen in these objects (see \\citealt{woodsthompson06,mereghetti08}, for review; see McGill SGR/AXP Online Catalog\\footnote{http://www.physics.mcgill.ca/$\\sim$pulsar/magnetar/main.html} for observational details). Two notable (and formerly defining) properties of magnetars are their high (as compared to that of other neutron stars of a similar age) observed X-ray luminosities $\\Lx$ in quiescence and their non-detection at radio wavelengths. The first suggests that heat generated from the decay of a strong magnetic field is the source of their bright X-ray luminosity \\citep{thompsonduncan96,heylkulkarni98,colpietal00,aguileraetal08} since $\\Lx$ is greater than that available from their rotational and thermal reservoirs. The second suggests that magnetars do not emit in radio. However, recent observations have brought these characteristics into question, in particular, the discovery of X-ray luminosities lower than spin-down luminosities (or rate of rotational energy loss $\\Edot$) for, and radio emission from, several magnetars (see \\citealt{reaetal12}, and references therein). The blurring of distinctions between magnetars and normal rotation-powered pulsars suggests magnetars may simply be a different manifestation of normal pulsars (see, e.g., \\citealt{kaspi10,pernapons11,ponsperna11}), although radio emission from magnetars do show some behavior that are different than radio emission from normal pulsars (see discussion in Sec.~2 of \\citealt{reaetal12}, and references therein).  In light of these latest discoveries, \\citet{reaetal12} examine the observed properties of radio active and inactive magnetars. They notice that all radio magnetars have X-ray efficiency $\\Lx/\\Edot<1$. They also calculate the electric potential difference across the magnetic pole $\\Delta\\Phi$ (or voltage gap) for magnetars and normal pulsars and find an apparent anti-correlation between voltage gap and X-ray efficiency for magnetars. They then conduct simulations which can produce an anti-correlation between $\\Delta\\Phi$ and $\\Lx/\\Edot$, depending on the neutron star magnetic field at birth. They conclude that there exists a fundamental plane ($\\Delta\\Phi$ versus $\\Lx/\\Edot$) for radio magnetars, in which a magnetar will be radio active if $\\Lx/\\Edot<1$ and the voltage gap is large and radio inactive if $\\Lx/\\Edot>1$.  Here we point out that one needs to be careful about claiming trends and correlations between parameters (e.g., $\\Delta\\Phi$ and $\\Lx/\\Edot$) when, in fact, these parameters are not entirely independent. We also extend the analysis of \\citet{reaetal12} by considering magnetars and other X-ray bright pulsars within the unified picture of neutron stars, as outlined in \\citet{kaspi10}.  In Section~\\ref{sec:model}, we briefly describe the standard model for pulsars and some of the basic equations derived from this model, as well as mention relevant recent works. In Section~\\ref{sec:results}, we show results from using these equations and the observed properties of magnetars and other pulsars. We summarize our results and discuss their implications in Section~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  We showed that using the two prime pulsar observables, spin period $P$ and period derivative $\\Pdot$, to calculate the voltage gap $\\Delta\\Phi$, as derived from standard pulsar theory \\citep{goldreichjulian69}, yields no new information beyond what can be inferred from the spin-down luminosity $\\Edot$. As a consequence, the parameter space or plane spanning $\\Delta\\Phi$ and $\\Lx/\\Edot$ is a simple transformation of, and no more fundamental or optimal than, the plane spanning $\\Edot$ and $\\Lx$. We showed that trends (in $\\Delta\\Phi$ and $\\Lx/\\Edot$ parameter space) seen among sources can be easily understood from standard pulsar theory and do not require complex simulations for explanation. For example, the anti-correlation between $\\Delta\\Phi$ and $\\Lx/\\Edot$ is a deception due to their particular dependence on $P$ and $\\Pdot$. Finally, we showed that a condition for pulsar radio activity/inactivity based on X-ray efficiency ($\\Lx/\\Edot\\sim 1$) seems to hold true in the unified picture of neutron stars, although there exists sufficient uncertainties for some sources that they could invalidate this conclusion. Thus the mechanism for pulsar radio emission appears to be similar among the different classes of neutron stars. However there are important differences between the observed properties of radio emission from magnetars and normal pulsars (see \\citealt{reaetal12}, and references therein), and these are likely due in part to differences in emission location and magnetic field strength/geometry of the magnetosphere (see, e.g., \\citealt{beloborodovthompson07,beloborodov12}, and references therein). New sources and improved measurements would provide a better understanding of radio behavior, as well as continuing theoretical work.  Up to this point, we have not discussed one other parameter, magnetic field $B$, that is often used to compare magnetars and pulsars. The most common method of determining $B$ for individual sources is by using eq.~(\\ref{eq:magb}) [or eq.~(\\ref{eq:magbspit}); see also \\citealt{glampedakisandersson11}, where it is argued that using eq.~(\\ref{eq:magb}) for magnetars leads to an overestimate of $B$]. In this case, our statements regarding comparisons between parameters that are only derived from the two observables $P$ and $\\Pdot$ apply here as well.  For example, plots of $B$ versus $\\tauc$ provide no additional information than what is contained within the standard pulsar $P$-$\\Pdot$ diagram, while plots of surface temperature $T$ versus $B$ (see, e.g., \\citealt{ponsetal07,zhuetal09}) and $T$ versus $\\tauc$ (see, e.g., \\citealt{kaplanvankerkwijk11}) are similar to $\\Lx$-$\\Edot$ (though there are systematic differences between measurements of $T$ and $\\Lx$).  On the other hand, independent measurements of magnetic field (from, e.g., spectral lines) or true ages (from, e.g., supernova remnants) do yield new information and are thus extremely valuable. In regards to the latter, we note that, for young pulsars, characteristic age generally disagrees with true age in cases where both can be determined (see, e.g., \\citealt{hoandersson12}).  Finally, the reason why magnetars exhibit high X-ray luminosities $\\Lx$ (for their age) is not known for certain. What is known is that an additional source of internal heat (beyond residual heat from neutron star formation) must be present in the outer crust \\citep{kaminkeretal06,kaminkeretal09,hoetal12}. Magnetic field evolution and decay could provide this heat source (see, e.g., \\citealt{ponsetal07,ponsetal09,cooperkaplan10,priceetal12}). We note that several magnetars may have very similar X-ray luminosities (see Fig.~\\ref{fig:edotlx}; see also \\citealt{durantvankerkwijk06}); this could suggest that their crustal field strengths are similar and the field decay timescale is longer than the age of the oldest of these sources."
    },
    "1208/1208.6538_arXiv.txt": {
        "abstract": "The Gamow-Teller response is astrophysically important for a number of nuclides, particularly around iron. The random phase approximation (RPA) is an efficient way to generate strength distributions. In order to better understand both theoretical systematics and uncertainties, we compare the Gamow-Teller strength distributions for a suite of nuclides and for a suite of interactions, including semi-realistic interactions in the $1p$-$0f$ space with the RPA and a separable multi-shell interaction in the quasi-particle RPA. We also compare with experimental results where available. ",
        "introduction": "Gamow-Teller (GT) electron capture ($\\beta$-decay) transitions, caused by the $\\sigma\\tau_+$ ($\\sigma\\tau_-$) operator, are some of the most important nuclear weak processes in astrophysics. For a review of spin-isospin transitions see Ref. \\cite{Ost92}. The GT transitions in $fp$-shell nuclei play important roles at the core collapse stages of supernovae, specially in neutrino induced processes. One of the factors controlling the gravitational core-collapse of massive stars is the lepton fraction; the lepton fraction in turn is governed by $\\beta$-decay and electron capture rates among iron-regime nuclides. A primary and non-trivial contribution to the weak rates is the distribution of GT strength. GT strengths have important implications in other astrophysical scenarios as well, such as explosive nucleosynthesis in O-Ne-Mg white dwarfs (see Ref. \\cite{NR07} and references therein) .  GT distributions have been extracted experimentally using different techniques. Whereas the $\\beta$-decay extraction is done in a model independent manner and are used to calibrate the B(GT), the charge-exchange reactions require further assumptions and the resulting extraction of GT distributions cannot be truly done in a model-independent manner. Consequently astrophysical calculations rely either upon crude estimates or upon more detailed microscopic calculations. The main difficulty with both experiment and theory is that the strength distribution connects to many states. Further in astrophysical environments one needs finite temperature GT strength functions as the temperature is high enough for excited states in the parent to be thermally populated.  The isovector response of nuclei may be studied using the nucleon charge-exchange reactions $(p,n)$ or $(n,p)$; by other reactions such as $(^{3}$He,$t$), ($d,^{2}$He) or through heavy ion reactions. The $0^{0}$ GT cross sections ($\\Delta T =1, \\Delta S =1, \\Delta L =0, \\hspace{0.1cm} 0\\hbar\\omega$ excitations) are proportional to the analogous beta-decay strengths. Charge-exchange reactions at small momentum transfer can therefore be used to study beta-decay strength distributions when beta-decay is not energetically possible. The $(p,n)$,  $(^{3}$He,$t)$ reactions probe the GT$_{-}$ strength (corresponding to $\\beta^{-}$-decay) and the $(n,p)$, $(d,^{2}$He) reactions give the strength for $\\beta^{+}$-decay/electron capture, i.e. GT$_{+}$ strength. The study of $(p,n)$ reactions has the advantage over $\\beta$-decay measurements in that the GT$_{-}$ strength can be investigated over a large region of excitation energy in the residual nucleus. On the other hand the $(n,p)$ reactions populates only $T = T_{0}+1$ states in all nuclei heavier than $^{3}$He. This means that other final states (including the isobaric analog resonance) are forbidden and GT$_{+}$ transitions can be observed relatively free of background. The study of these reactions suggest that a reduction in the amount of GT strength is observed relative to theoretical calculations. The GT quenching is on the order of 30-40 $\\%$ \\cite{Vet89}.  Theory for GT transitions falls generally into three camps: simple independent-particle models (e.g.  Ref. \\cite{Ful80}); full-scale interacting shell-model calculations; and, in between, the random-phase approximation (RPA) and quasi-particle random-phase approximation (QRPA). Independent-particle models underestimate the the total GT strength, because the Fermi surface is insufficiently fragmented, while  also placing the centroid of the GT strength too high for  even-even parent nuclides and too low  on odd-A and odd-odd parents \\cite{Pin00}. Full interacting shell-model calculations are computationally demanding, although one can exploit the Lanczos algorithm, commonly used in large shell-model diagonalization \\cite{Whi72}, to efficiently generate the strength distribution \\cite{SMreview05}; for medium-mass nuclei one still needs to choose from among a number of competing semi-realistic/semi-empirical interactions. RPA and QRPA can be thought of as approximations to a full shell-model calculation and are much less demanding computationally.  In this paper we compare GT strength distributions for a suite of iron-region nuclides relevant to astrophysics: $^{54,55,56}$Fe,  and $^{56,58}$Ni. For each of these nuclides we compute the GT strengths in several RPA calculations. Each RPA calculation is in occupation rather than configuration space, which is appropriate inasmuch as the GT operator only affects spin and isospin. (In fact each calculation properly speaking is proton-neutron RPA or QRPA, as the RPA/QRPA phonon operators change protons into neutrons or vice-versa.) The calculations, which will be described in greater detail in subsequent sections, are :  \\noindent $\\bullet$ pn-RPA in a major harmonic oscillator shell, that is, the $1p$-$0f$ shell, with three different semi-realistic/semi-empirical interactions \\cite{Pov01,Ric91,Hon02}. For details see Section 2.  \\noindent $\\bullet$ pn-QRPA in a multi-shell single-particle space with a schematic interaction that has been previously applied to similar calculations \\cite{Mut92}. For further details we refer to Section 3.   These particular calculations were chosen because of the availability of codes; one could imagine a larger set of calculations (e.g., pn-QRPA with semi-realistic interactions) but relevant codes either do not exist or are not available to us.  These calculations will help to understand systematic similarities and differences between (a) different $1p$-$0f$ shell-model interactions and (b) between $0\\hbar\\Omega$ shell-model calculations against multi-shell calculations with a separable interaction. For example, for some cases in the $1s$-$0d$ shell the separable interactions yield a larger total strength and a higher centroid \\cite{Nab99} than shell-model calculations.  Section~4 presents the results and discussions on use of various pn-RPA schemes. We finally present the summary and conclusions in Section~5. ",
        "conclusions": "Under astrophysical conditions, both the electron capture and beta decay of fp-shell nuclei depend heavily on the centroid placement and total strength of the calculated Gamow-Teller strength distributions. In this work we  presented a comparative study of the Gamow-Teller strength distributions  for a suite of astrophyiscally important fp-shell nuclide ($^{54,55,56}$Fe,  and $^{56,58}$Ni) using a suite of interactions, including semi-realistic interactions in the $1p$-$0f$ space with the RPA and a separable multi-shell interaction in the quasi-particle RPA. Where possible, we also presented comparison with measured data. We further compared and contrasted the statistics of calculated and measured GT strength functions using various pn-RPA schemes in this paper. Our calculations satisfied the model independent Ikeda sum rule. Work is currently in progress for other important odd-A and odd-odd cases.  The QRPA model places the centroid at much lower energies in daughter nuclei as compared to other RPA interactions. Further the placement of GT centroids by the pn-QRPA model is, in general, in good agreement with the centroids of the measured data. This tendency of QRPA model can favor higher values of electron capture rates in stellar environment and can bear significance for astrophysical problems. On the other extreme, the GXPF1 interaction usually leads to placement of GT centroid at much higher energies in daughter compared to other pn-RPA interactions.  The present study showed that the total strengths, using various RPA interactions, were in better agreement with the measured data when compared to the QRPA calculated strength. Further the width of the strength functions calculated within the QRPA was much larger than those calculated with other RPA interactions. For the special $N = Z$ nucleus $^{56}$Ni, the QRPA model calculated Gamow-Teller strength function was well fragmented as compared to other RPA interactions (including the recently used GXPF1J interaction)  and may lead to interesting consequences for heavy element nucleosynthesis.   \\vspace{0.5in} \\textbf"
    },
    "1208/1208.6012_arXiv.txt": {
        "abstract": "{The source of Ultra High Energy Cosmic Rays (UHECR) is still an unresolved mystery. Up until recently, sources of Gamma Ray Bursts (GRBs) had been considered as a suitable source for UHECR. Within the fireball model, the UHECR produced at GRBs should be accompanied with a neutrino flux detectable at the neutrino telescope such as IceCube. Recently, IceCube has set an upper bound on the neutrino flux accompanied by GRBs about 3.7 times below the prediction. We investigate whether this deficit can be explained by the oscillation of the active neutrinos to sterile neutrinos {\\it en route} from the source to the detectors within the pseudo-Dirac scenario. We then discuss the implication of this scenario for diffuse supernova relic neutrinos.}   \\begin{document} ",
        "introduction": "}    One of the challenges in cosmic ray physics is identifying the source of Ultra High Energy Cosmic Rays (UHECRs) with energy $>10^{18}$~eV. Simple calculation of the energy budget and rate of  astrophysical explosions shows that the transient Gamma Ray Bursts (GRBs)  can accommodate the observed flux of UHECRs~\\cite{Waxman:1995vg}. The GRBs are energetic explosions with luminosity $\\sim10^{51}$~erg/s at cosmological distances which have been observed by various satellites via their gamma ray emissions at a  rate $\\sim$~2/day.    Within the fireball model of the GRB explosion ({\\it see}~\\cite{Waxman:2003vh} for a review on this model), the decay of the photo-pions produced by accelerated protons in the expanding shock wave generates a flux of accompanying neutrinos. Based on the observed flux of UHECRs, the expected flux of neutrinos is $E_\\nu^2dN_\\nu/dE_\\nu\\sim 5\\times 10^{-9}\\,{\\rm GeV}{\\rm cm}^{-2}{\\rm s}^{-1}{\\rm sr}^{-1}$ in the energy range $\\sim 100\\,{\\rm TeV}-10\\,{\\rm PeV}$ (the so-called Waxman-Bahcall flux~\\cite{Waxman:1997ti}). On the other hand within the fireball model, the neutrino flux emitted from each GRB can be calculated from the observation of gamma ray energy spectrum. Normalizations based on the cosmic ray flux and the gamma ray flux both result in approximately the same prediction for the neutrino flux. The Waxman-Bahcall flux leads to $\\sim$ 10 events per year in a km$^3$-scale neutrino telescope.   Recently, the IceCube experiment published the result of the analysis of data taken during the construction of detector, that is IC40 and IC59 \\cite{Abbasi}. The combined analysis of the data does not show any neutrino signal correlated with the observed GRBs during the data-taking period. The limit of IceCube on the neutrino flux is $\\sim 3.7$ times smaller than the prediction \\cite{Abbasi}. From this limit, the IceCube collaboration has concluded that the GRBs cannot account for the bulk of the UHECRs in the universe \\cite{Abbasi}. However, this conclusion can be questioned in three ways: (1) As shown in \\cite{Winter}, within the fireball model, uncertainties are so large that the prediction for the neutrino flux can be lower than the standard Waxman-Bahcall prediction by up to a factor of ten (see also~\\cite{Murase:2005hy}). (2) In \\cite{Dar}, it is discussed that within the cannonball model, the sources of GRBs can account for the UHECR without violating the neutrino flux bound because the cannonball model does not predict a sizable neutrino flux. (3) Finally, it is possible that some beyond standard model effects reduce the flux after production (see {\\it e.g.} \\cite{Barranco}). In this letter, we entertain the third possibility.   Although overwhelming evidence has been gathered showing that there are at least two massive neutrino mass eigenstates, we do not still know whether neutrinos are of Majorana type or of Dirac type. An interesting possibility is to have pseudo-Dirac neutrino mass scheme~\\cite{Wolfenstein:1981kw}, where the dominant contribution to the neutrino mass comes from the Dirac mass term, $m_D$, with a small correction from the Majorana mass term, $m_M$ such that $m_M\\ll m_D$. Notice that the presence of the Dirac term requires right-handed sterile neutrinos. Within the pseudo-Dirac scenario, there is a small mass splitting between active and sterile neutrinos making the oscillation between active and sterile components in principle possible. For very small $m_M$, the oscillation length will be too large for the atmospheric and solar neutrinos to oscillate to sterile ones. However for neutrinos traveling over cosmological distances, the baseline will be large enough for active-sterile oscillation to take place~\\cite{pD}. In this paper, we investigate this possibility to relax the conflict between IceCube bound on the neutrinos from GRBs and the expected neutrino flux from GRBs as the origin of UHECR.   The paper is organized as follows: In section~\\ref{sec:scenario}, we briefly discuss the pseudo-Dirac scenario for neutrino masses. In section~\\ref{sec:grb}, we investigate the possible depletion of GRB neutrino flux due to the presence of tiny active-sterile mass-squared difference predicted in pseudo-Dirac scenario. Section~\\ref{sec:sn} is devoted to the implications of pseudo-Dirac scenario for supernova neutrinos, discussing both single source and diffuse fluxes. The conclusions are summarized in section~\\ref{sec:conc}. ",
        "conclusions": ""
    },
    "1208/1208.5227_arXiv.txt": {
        "abstract": "We present a method for calculating the maximum elastic quadrupolar deformations of relativistic stars, generalizing the previous Newtonian, Cowling approximation integral given by [G.\\ Ushomirsky \\emph{et al.}, Mon.\\ Not.\\ R.\\ Astron.\\ Soc.\\ {\\bf{319}}, 902 (2000)]. (We also present a method for Newtonian gravity with no Cowling approximation.) We apply these methods to the $m=2$ quadrupoles most relevant for gravitational radiation in three cases: crustal deformations, deformations of crystalline cores of hadron--quark hybrid stars, and deformations of entirely crystalline color superconducting quark stars. In all cases, we find suppressions of the quadrupole due to relativity compared to the Newtonian Cowling approximation, particularly for compact stars. For the crust these suppressions are up to a factor of $\\sim 6$, for hybrid stars they are up to $\\sim 4$, and for solid quark stars they are at most $\\sim 2$, with slight enhancements instead for low mass stars. We also explore ranges of masses and equations of state more than in previous work, and find that for some parameters the maximum quadrupoles can still be very large. Even with the relativistic suppressions, we find that $1.4\\Msolar$ stars can sustain crustal quadrupoles of $\\text{a few}\\times10^{39}\\,\\text{g cm}^2$ for the SLy equation of state, or close to $10^{40}\\,\\text{g cm}^2$ for equations of state that produce less compact stars. Solid quark stars of $1.4\\Msolar$ can sustain quadrupoles of around $10^{44}\\,\\text{g cm}^2$. Hybrid stars typically do not have solid cores at $1.4\\Msolar$, but the most massive ones ($\\sim 2\\Msolar$) can sustain quadrupoles of $\\text{a few}\\times10^{41}\\,\\text{g cm}^2$ for typical microphysical parameters and $\\text{a few}\\times10^{42}\\,\\text{g cm}^2$ for extreme ones. All of these quadrupoles assume a breaking strain of $10^{-1}$ and can be divided by $10^{45}\\,\\text{g cm}^2$ to yield the fiducial ``ellipticities'' quoted elsewhere. ",
        "introduction": "Shortly after the discovery of pulsars and the realization that they are rotating neutron stars, deformations of rotating neutron stars were proposed as sources of continuous gravitational radiation~\\cite{Shklovskii1969, Ostriker1969, Ferrari1969, Melosh1969}; see~\\cite{Press1972} for an early review. Searches for such radiation are an ongoing concern of the LIGO and Virgo gravitational wave detectors~\\cite{LIGO_psrs2010, LIGO_CasA, LIGO_Vela}; see~\\cite{Owen2009, Pitkin, Astone} for recent reviews. It is thus of great interest to know the maximum quadrupolar deformation that a neutron star could sustain, in order to motivate further searches and help interpret upper limits or detections. In the case of elastic (as opposed to magnetic) deformations, the main factor influencing the answer is whether the neutron star contains particles more exotic than neutrons~\\cite{OwenPRL, Owen2009}. However, the structure of the star also plays an important role.  While there are relativistic calculations of the quadrupole deformations due to magnetic fields (e.g.,~\\cite{IS,CFG,YKS, FR,CR}), all the computations involving elastic deformations have used Newtonian gravity. Moreover, all but two of these computations have used the integral expression obtained in the Cowling approximation (i.e., neglecting the self-gravity of the perturbation) by Ushomirsky, Cutler, and Bildsten (UCB)~\\cite{UCB}; see \\cite{OwenPRL, Lin, KS, Horowitz}. Haskell, Jones, and Andersson (HJA)~\\cite{HJA} dropped the Cowling approximation using a somewhat different formalism than UCB's; there is a further application of their results in~\\cite{Haskelletal}.  We improve these treatments by generalizing the UCB integral to relativistic gravity with no Cowling approximation. We also provide a similar generalization for the Newtonian no-Cowling case, as a warm-up. In addition to providing a simpler formalism for performing computations than the more general Newtonian gravity treatment in HJA, the integrals we obtain allow us to verify that a maximal uniform strain continues to yield the maximum quadrupole deformation in the Newtonian and relativistic no-Cowling cases. (UCB showed this to be true for an arbitrary equation of state in the Newtonian Cowling approximation case; we are able to verify that it is true in the more general cases for each background stellar model we consider.)  We then apply our calculation to the standard case of quadrupoles supported by shearing the lattice of nuclei in the crust, as well as the cases where the quadrupole is supported by the hadron--quark mixed phase lattice in the core, or a crystalline color superconducting phase throughout a solid strange quark star. For the crustal quadrupoles, we calculate the shear modulus following HJA, using the equation of state (EOS) and composition results of Douchin and Haensel~\\cite{DH} and the effective shear modulus calculated by Ogata and Ichimaru~\\cite{OI}. (There are recent improvements to the Ogata and Ichimaru result~\\cite{HH, Baiko,BaikoCPP}, but these only reduce their shear modulus by $<10\\%$.) For the hadron--quark mixed phase, we use our recent calculations of the EOS and shear modulus~\\cite{J-MO1} for a variety of parameters. (We also consider the range of surface tensions for which the mixed phase is favored.) For crystalline quark matter, we use the shear modulus calculated by Mannarelli, Rajagopal, and Sharma~\\cite{MRS}, and the EOS given by Kurkela, Romatschke, and Vuorinen~\\cite{KRV}.  In all cases, we use a breaking strain of $0.1$, comparable to that calculated by Horowitz and Kadau~\\cite{HK} using molecular dynamics simulations. (Hoffman and Heyl~\\cite{HoHe} have recently obtained very similar values over more of parameter space.) This result is directly applicable to the crustal lattice, at least for the outer crust, above neutron drip (though see Chugunov and Horowitz~\\cite{ChHo} for caveats). We also feel justified in applying it to the inner crust, as well as to the mixed phase and crystalline quark matter, since the primary source of the high breaking strain appears to be the system's large pressure. But one can apply our results to any breaking strain using the linear scaling of the maximum quadrupole with breaking strain.  In our general relativistic calculation, we use the relativistic theory of elasticity given by Carter and Quintana~\\cite{CQ} and placed in a more modern guise by Karlovini and Samuelsson~\\cite{KaSaI}. However, all we need from it is the relativistic form of the elastic stress-energy tensor, which can be obtained by simple covariance arguments, as noted by Schumaker and Thorne~\\cite{ST}. We also use the standard \\citet{TC} Regge-Wheeler gauge~\\cite{RW} formalism for perturbations of static relativistic stars, following Hinderer's recent calculation~\\cite{Hinderer} of the quadrupole moment of a tidally deformed relativistic star (first discussed in~\\citet{FH}), and the classic calculation by \\citet{Ipser}.  Even though we are interested in the gravitational radiation emitted by rotating stars, it is sufficient for us to calculate the static quadrupole deformation. As discussed by Ipser~\\cite{Ipser}, and then proved for more general situations by Thorne~\\cite{Thorne}, this static quadrupole (obtained from the asymptotic form of the metric) can be inserted into the quadrupole formula to obtain the emitted gravitational radiation in the fully relativistic, slow-motion limit. [This approximation has uncontrolled remainders of order $(\\omega/\\omega_K)^2$, where $\\omega$ and $\\omega_K$ are the star's angular velocity and its maximum---i.e., Kepler---angular velocity, respectively. This ratio is $\\lesssim 10^{-2}$ for the pulsars for which LIGO has been able to beat the spin-down limit~\\cite{LIGO_psrs2010}.]  We shall generally show the gravitational constant $G$ and speed of light $c$ explicitly, though we shall take $G = c =1$ in most of Sec.~\\ref{GR}, only restoring them in our final expressions. The relativistic calculation was aided by use of the computer algebra system {\\sc{Maple}} and the associated tensor manipulation package {\\sc{GRTensorII}}~\\cite{GRTensor}. We used {\\sc{Mathematica}}~7 to perform numerical computations.  The paper is structured as follows: In Sec.~\\ref{Newt}, we review UCB's formalism and extend it by introducing a Green function to compute the maximum Newtonian quadrupole deformation without making the Cowling approximation. In Sec.~\\ref{GR}, we further generalize to the fully relativistic case, and compare the various approximations for the maximum quadrupole. In Sec.~\\ref{results2}, we show the maximum quadrupoles for three different cases: first crustal quadrupoles, then hadron--quark hybrid quadrupoles, and finally solid strange quark star quadrupoles. We also describe the modifications to our formalism needed to treat solid strange quark stars. We discuss all these results in Sec.~\\ref{discussion}, and summarize and conclude in Sec.~\\ref{concl2}. In the Appendix, we show that the mixed phase is favored by global energy arguments even for surface tensions large enough that it is disfavored by local energy arguments. ",
        "conclusions": "\\label{concl2}  We have presented a method for calculating the maximum elastic quadrupole deformation of a relativistic star with a known shear modulus and breaking strain. We then applied this method to stars whose elastic deformations are supported by a shear modulus either from the Coulomb lattice of nuclei in the crust, a hadron--quark mixed phase in the core, or crystalline superconducting strange quark matter throughout the star. (In the last case, we have made the requisite changes to the method so that it is valid when the star has a nonzero surface density and the shear modulus does not vanish at the star's surface.) In all but the strange quark case, we find that the relativistic quadrupole is suppressed, compared with the standard, Newtonian Cowling approximation calculation of the quadrupole, at least for stars with masses of $\\gtrsim 1\\Msolar$ (corresponding to the observed masses of neutron stars) and the EOSs we have investigated. These suppressions can be up to $\\sim 4$ in the hybrid case, and $\\sim 6$ in the crustal case. In the strange quark star case, the Newtonian Cowling approximation calculation slightly underestimates the quadrupole (by tens of percent) for low-to-standard mass stars, but is still an overestimate of $\\sim 2$ at higher masses.  These suppressions strengthen the Horowitz~\\cite{Horowitz} argument that searches for gravitational waves from elastically deformed neutron stars supported by crustal stresses are biased towards lower-mass stars. The same argument also applies to strange quark stars, though there the suppressions with increasing mass are less severe (and the maximum quadrupoles are all considerably larger). However, this argument does not apply to quadrupole deformations of hybrid stars, since the increase in the size of the region of mixed phase with increasing mass dominates the various suppressions.  Our results also imply that many of the previous calculations of elastic quadrupoles (e.g.,~\\cite{Lin, Haskelletal, KS, UCB, HJA}) will need their results revised downwards. (While we find much larger maximum quadrupoles for solid strange quark stars than did Lin~\\cite{Lin}, this is only because we assume a breaking strain $10$ times that assumed by Lin. If we take the same $10^{-2}$ breaking strain as does Lin, then we find a suppression of a factor of a few, though this is very likely within the uncertainties of Lin's calculation, which assumed a uniform density, incompressible star with a uniform shear modulus.)  It is instructive to compare our results with the numbers quoted in Pitkin's review~\\cite{Pitkin}. All of these were obtained by Pitkin using scalings given in the aforementioned papers, sometimes updating to the Horowitz and Kadau~\\cite{HK} breaking strain, and provide a good overview of the standard Newtonian predictions. None of our detailed calculations for maximum crustal quadrupoles approach the high values Pitkin obtained using UCB's fitting formula (as corrected by \\citet{OwenPRL}). However, our very largest hybrid star quadrupoles are an order of magnitude above Pitkin's quoted maximum, even if one only assumes a breaking strain of $10^{-2}$, as does Pitkin. Additionally, our estimates for maximum solid quark star quadrupoles ($\\sim 10^{44}\\text{ g cm}^2$ for $1.4\\Msolar$ stars) are considerably larger than the ones quoted by Pitkin (based on a different shear modulus model), even if we reduce them by an order of magnitude due to scaling the breaking strain to Pitkin's $10^{-2}$. In fact, they are in the same range as those Pitkin quotes for a model for crystalline superconducting hybrid stars (with an optimistic gap parameter $5$ times the one we used for solid quark stars, leading to a shear modulus $\\sim 40$ times our shear modulus's maximum value).  Even with the relativistic suppressions, we obtain maximum quadrupole deformations of $\\text{a few}\\times10^{42}\\text{ g cm}^2$ in the hybrid case for a very stiff hadronic EOS, and $\\text{a few}\\times10^{41}\\text{ g cm}^2$ for more realistic cases. In both situations, the largest maximum quadrupoles are given by the most massive stars. These values are proportional to the breaking strain and assume that the Horowitz and Kadau~\\cite{HK} breaking strain of about $0.1$ is applicable to the mixed phase. Such large quadrupole deformations were previously thought only to be possible for solid quark stars (see~\\cite{OwenPRL, Lin, Haskelletal, KS}), or from crustal deformations in the very low-mass neutron stars considered by Horowitz~\\cite{Horowitz}.  These large deformations (corresponding to fiducial ellipticities of $\\text{a few}\\times10^{-3}$ in the extreme case, and $\\sim 5\\times10^{-4}$ in a more realistic case) would be able to be detected by current LIGO searches for gravitational waves from certain known neutron stars~\\cite{LIGO_psrs2010, LIGO_CasA, LIGO_Vela}. (However, we must note that there is no reason to assume that such isolated stars are anywhere near maximally strained, even neglecting the uncertainties in the description of their interiors.)  The prospects for crustal quadrupoles are now somewhat less optimistic, and definitely favor lower-mass stars. However, for a canonical $1.4\\Msolar$ neutron star, we find that the maximum relativistic crustal quadrupole is in the range $\\sim\\text{(1--6)}\\times 10^{39}\\text{ g cm}^2$ [corresponding to fiducial ellipticities of $\\sim\\text{(1--8)}\\times 10^{-6}$], depending on the model used for the crust and the high-density EOS. (Note that the fully consistent Douchin and Haensel model with its associated high-density EOS yields the lowest numbers. Additionally, there is the possibility of a further reduction of up to $\\sim 2$ due to the angle averaging procedure used to obtain the shear modulus.) On the high side, these numbers are consistent with those given previously for breaking strains of $0.1$ by Horowitz~\\cite{HK, Horowitz},\\footnote{But recall that the results from Horowitz~\\cite{Horowitz} were obtained using the SLy EOS and crustal composition results, so they are the same as our Newtonian Cowling approximation SLy predictions, given in Fig.~\\ref{Q22s_vs_M_SLy_DH_HJA}, except $\\sim7\\%$ lower, since Horowitz is using the Horowitz and Hughto~\\cite{HH} result for the shear modulus. In the fully relativistic case, one requires a thicker crust than provided by the pure SLy results to obtain values for the maximum quadrupole comparable to those given by Horowitz.} though they are a factor of $\\sim 5$ lower than the maximum Pitkin~\\cite{Pitkin} obtained using scalings of previous results and the maximum value given by HJA (scaled to this breaking strain). For stars around $2\\Msolar$, the relativistic suppressions lead to maximum quadrupoles that are nearly an order of magnitude smaller than those for a $1.4\\Msolar$ star in the compact SLy case: ${\\sim\\text{(1--5)}}\\times 10^{38}\\text{ g cm}^2$ [corresponding to fiducial ellipticities of $\\sim\\text{(1--6)}\\times 10^{-7}$]; and even in the much less compact LKR$1$ case, there is a suppression of $\\sim 5$. Previous Newtonian studies (see Fig.~3 in~\\cite{Horowitz}) had only found suppressions of around a factor of $4$, due to the thinning of the crust and the increase in Newtonian gravity with increasing mass. It will be interesting to consider further models for the crustal composition and EOS in this case, particularly the large suite of crustal models including the pasta phases recently calculated by Newton, Gearheart, and Li~\\cite{NGL}. (See~\\cite{GNHL} for order-of-magnitude estimates of the maximum quadrupole for these models, illustrating the sensitive dependence on the slope of the symmetry energy.)  One can also compare these maximum elastic quadrupoles with those generated by an internal magnetic field. Here the values depend, of course, upon the equation of state, compactness, and---perhaps most crucially---magnetic field topology, as well as the quantity one chooses to use to measure the magnitude of the magnetic field. But sticking to order-of-magnitude numbers, and considering a canonical $1.4\\Msolar$ neutron star, Frieben and Rezzolla~\\cite{FR} show that a toroidal internal field of $\\sim10^{15}$~G would generate a quadrupole of $\\sim 10^{39}$--$10^{40}\\text{ g cm}^2$, comparable to the maxima we find for crustal quadrupoles. Similarly, quadrupoles of $\\sim 10^{41}$--$10^{42}\\text{ g cm}^2$, around the maxima we find for hybrid stars, could come from magnetic fields of $\\sim 10^{16}$~G, while the maximum quadrupoles of $\\sim 10^{44}\\text{ g cm}^2$ we find for crystalline strange quark stars could also be generated by magnetic fields of $\\sim 10^{17}$~G, close to the maximum allowed field strength. (But note that these magnetic deformations are all computed for ordinary, purely hadronic neutron stars. Additionally, the quoted maximum elastic quadrupoles in the hybrid case are attained only for more massive stars than the $1.4\\Msolar$ stars for which we are quoting the magnetic deformation results.) The quoted values for magnetic quadrupoles come from the fits given in Sec.~7 of Frieben and Rezzolla~\\cite{FR}, except for the final ones, which are obtained from inspection of their Fig.~5 and Table~3. All these values agree in order of magnitude with the predictions for the twisted torus topology given by Ciolfi, Ferrari, and Gualtieri~\\cite{CFG}, and with many other studies for various topologies cited in Frieben and Rezzolla~\\cite{FR}. But note that very recent calculations by Ciolfi and Rezzolla~\\cite{CR} show that the magnetic field required to obtain a given quadrupole deformation with the twisted torus topology could be reduced by about an order of magnitude if the toroidal contribution dominates.  One would also like to make relativistic calculations of the maximum energy that could be stored in an elastic deformation. This would be useful in properly computing the available energy for magnetar flares, for instance. (Using Newtonian scalings,  \\citet{CO} estimated that the hybrid case was especially interesting compared to existing LIGO upper limits for gravitational wave emission from such flares.) The basic expressions (at least in the perfect fluid case) appear to be readily available in the literature (see, e.g.,~\\cite{Schutz2,DI}; \\cite{ST, Finn} give related results including elasticity). However, one cannot apply these directly to the crustal and hybrid cases, even in the Newtonian limit, due to the distributional nature of the density and pressure perturbations. Specifically, the sudden change in shear modulus at the phase transitions gives delta functions in the derivatives of the density and pressure perturbations. Since the energy expressions involve squares of these derivatives, one would have to invoke some sort of regularization procedure, or apply a different method. Developing appropriate expressions for this case will be the subject of future work.  Returning to the quadrupoles, one might also want to consider the shape of the deformed star, particularly in the relativistic case---the ellipticity is already only a rough indicator of the shape of the deformation in the Newtonian case---as has now been done in~\\cite{NKJ-M_shape}. But the effects of the star's magnetic field are surely the most interesting to consider, from its influence on the lattices that support elastic deformations, to the changes to the boundary conditions at the star's surface from an external magnetic field (particularly for magnetars), to the internal magnetic field's own contribution to the star's deformation. One might also want to consider the lattice's full elastic modulus tensor in this case, instead of simply assuming a polycrystalline structure and angle averaging to obtain an effective isotropic shear modulus, as was done here. (And even if one assumes a polycrystalline structure, one could use more involved, sharper bounds on the shear modulus than the ones considered here---see~\\cite{WDO} for a classic review of such bounds.)"
    },
    "1208/1208.5682_arXiv.txt": {
        "abstract": "{\\emph{Context.} Owing to their computational simplicity, models with elliptical potentials (pseudo-elliptical) are often used in gravitational lensing applications, in particular for mass modeling using arcs and for arc statistics.  However, these models generally lead to negative mass distributions in some regions and to dumbbell-shaped surface density contours for high ellipticities.  \\emph{Aims.} We revisit the physical limitations of the pseudo-elliptical Navarro--Frenk--White (PNFW) model, focusing on the behavior of the mass distribution close to the tangential critical curve, where tangential arcs are expected to be formed.  We investigate the shape of the mass distribution on this region and the presence of negative convergence. We obtain a mapping from the PNFW to the NFW model with elliptical mass distribution (ENFW). We compare the arc cross section for both models, aiming to determine a domain of validity for the PNFW model in terms of its mass distribution and for the cross section.  \\emph{Method.} We defined a figure of merit to {\\it i}) measure the deviation of the iso-convergence contours of the PNFW model to an elliptical shape, {\\it ii}) assigned an ellipticity $\\varepsilon_\\Sigma$ to these contours, {\\it iii}) defined a corresponding iso-convergence contour for the ENFW model. We computed the arc cross section using the ``infinitesimal circular source approximation''.  \\emph{Results.} We extend previous work by investigating the shape of the mass distribution of the PNFW model for a broad range of the potential  ellipticity parameter $\\varepsilon$ and characteristic convergence $\\kappa_s^\\varphi$. We show that the maximum value of  $\\varepsilon$ to avoid dumbbell-shaped mass distributions is explicitly dependent on $\\kappa_s^\\varphi$, with higher ellipticities  ($\\varepsilon \\simeq 0.5$, i.e., $\\varepsilon_\\Sigma \\simeq 0.65$) allowed for small  $\\kappa_s^\\varphi$. We determine a relation between the ellipticity of the mass distribution $\\varepsilon_\\Sigma$ and $\\varepsilon$ valid for any ellipticity. We also derive the relation of characteristic convergences, obtaining a complete mapping from PNFW to ENFW models, and provide fitting formulae for connecting the parameters of both models. Using this mapping, the cross sections for both models are compared, setting additional constraints on the parameter space of the PNFW model such that it reproduces the ENFW results.  We also find that the negative convergence regions occur far from the arc formation region and should therefore not be a problem for studies with gravitational arcs.  \\emph{Conclusions.} We conclude that the PNFW model is well-suited to model an elliptical mass distribution on a larger $\\varepsilon$--$\\kappa_s^\\varphi$  parameter space than previously expected.   However, if we require the PNFW model to reproduce the arc cross section of the ENFW well, the ellipticity is more restricted, particularly for low $\\kappa_s^\\varphi$. The determination of a domain of validity for the PNFW model and the mapping to ENFW models could have implications for the use of PNFW models for the inverse modeling of lenses and for fast arc simulations, for example.} ",
        "introduction": "}  Gravitational arcs are powerful probes of  the mass distribution in galaxies \\citep{Koopmans09,Barnabe11,Suyu12} and galaxy clusters \\citep{kovner89,miralda93,hattori97} and can be used to constrain cosmological models \\citep{bartelmann1998,oguri01, golse2002, bartelmann03,2010Sci...329..924J}. The main techniques employed to extract information from gravitational arcs have been {\\it arc statistics} \\citep[i.e. counting the number of arcs in  lens samples,][]{ wu93,1994ApJ...431...74G,bartelmann94} and {\\it inverse modeling} \\citep[i.e. ``deprojecting'' the arcs in individual lens systems to determine the lens and the source,][]{1993A&A...273..367K,gravlens,golse2002,lensview,jullo07,2010Sci...329..924J}. The first requires large samples of arcs while the second needs detailed information on the lensing systems, typically imaging from space and lens and source redshifts.  These applications have triggered arc searches in surveys covering large areas \\citep{rscs, sdss1,legacy,CS82,more11,CASSOWARYmethod, SBAS7},  in large spectrocospic surveys \\citep{SLACSV,BELLSI}, and in surveys targeting known clusters \\citep{luppino, lcdcs,2005MNRAS.359..417S,sdss2, sogras10,2010MNRAS.404..325R,kausch2010, sogras}. Upcoming wide field imaging surveys, such as the Dark Energy Survey \\citep{des05,annis05} and the Large Synoptic Survey Telescope \\citep{lsst,lsst1}, will lead to the identification of larger samples of arcs in thousands of galaxies and galaxy clusters, well suited for arc statistics. Moreover, deep observations from space combined with massive spectroscopy have been obtained for a limited number of clusters and were used for detailed mass modeling \\citep[see, e.g., ][]{2005ApJ...621...53B,2010MNRAS.402L..44R,2010Sci...329..924J}.  The simplest models that can account for some observed properties of arcs (multiplicity, relative positions, morphology) are built from axial models by introducing an ellipticity either on the mass distribution  \\citep[elliptical models,][]{schramm90,barkana98,gravlens,oguri03} or on the lensing potential \\citep[pseudo-ellitpical models,][]{kochanek89,kassiola93,kneib01}. Most parametric analyses of arcs, both for the inverse modeling \\citep{jullo07} and for arc statistics \\citep{oguri02, oguri03}, involve one or more elliptical/pseudo-elliptical models (adding, in some cases, external shear and substructures).  Elliptical models, whose surface density is constant over ellipses, are more realistic than pseudo-elliptical ones. For example, elliptical models are motivated by the results of N-body simulations, which show that dark matter halos are triaxial \\citep{2002ApJ...574..538J,2007MNRAS.378...55M}, such that their overall mass distribution can be modeled at first order by ellipsoids, whose surface density contours are elliptical. In contrast, the surface density of pseudo-elliptical models  generally has a pathological behavior, exhibiting regions where it takes negative values \\citep{kochanek89} and presenting a ``dumbbell'' (or ``peanut'') shape for high ellipticities \\citep{kovner89, SEF,kassiola93}, which does not represent the mass distribution of most physical systems.  On the other hand, pseudo-elliptical models provide simple analytic solutions for some lensing quantities,  allowing for fast numerical methods to be implemented, whereas elliptical models require the computation of integrals, which are more demanding numerically \\citep{schramm90,gravlens}. Therefore, studies that require numerous evaluations of the lensing quantities often employ pseudo-elliptical models. For example, several studies using arc simulations have used these models \\citep{oguri02, mene03,mene07}. Popular codes for lens inversion \\citep{golse2002,lensview, jullo07} are implemented using this type of model, too.  It is therefore relevant to determine a ``domain of validity'' for pseudo-elliptical models such that the negative convergence appears far from the arc-forming region and the shape of their mass distribution is closer to elliptical. The determination of  such a validity region could be useful, for example, to evaluate if a set of model parameters derived from the inversion of a system with arcs  using the PNFW model is physically acceptable.  Even for methods that use the ENFW for lens inversion \\citep{gravlens,Suyu12}, the PNFW could be useful, within its domain of validity, for a faster coarser probing of the parameter space that would subsequently need to be refined with the elliptical model.  Once a pseudo-elliptical model is found to be acceptable, it is nevertheless necessary to provide a correspondence to  an elliptical model. This would be necessary, for example, to compare results derived from the mass modeling using observed arcs with theoretical predictions. We therefore need to establish a mapping among the parameters of the two models. This mapping could also be used to replace an elliptic model by its corresponding pseudo-elliptic in lensing simulations.  In this paper we  focus on the widely used Pseudo-Elliptical Navarro--Frenk--White (hereafter PNFW) model and investigate its domain of validity as well as the mapping to  the corresponding elliptical model  (hereafter ENFW). To test the equivalence among the two models we compute the cross section for arc formation, from which we obtain additional constraints on the model parameters.  The outline of this paper is as follows: In Sect.~\\ref{PNFW} we briefly review the PNFW lens model, introducing the conventions and parameter ranges to be used throughout this work. In Sect.~\\ref{phys_lim_pnfw} we discuss the region of arc formation and derive physical limits of the PNFW mass distribution in this region.  In Sect.~\\ref{mapping} we consider two ways for assigning an ellipticity to the PNFW surface density. In Sect.~\\ref{charac_convg_relation} we obtain a mapping from the PNFW to the ENFW models. In Sect.~\\ref{compar_sc} we compare the arc cross sections of the two models using the mapping relations. In Sect.~\\ref{s&c} we present the summary and concluding remarks. In Appendix~\\ref{ap1} we present useful relations to derive some lensing functions for pseudo-elliptical models. In Appendix~\\ref{useful-fit} we provide fitting formulae for the limits on the mass distribution of the PNFW model and for the mapping to the ENFW. ",
        "conclusions": ""
    },
    "1208/1208.5966_arXiv.txt": {
        "abstract": "{Shock fronts in young supernova remnants are the best candidates for being sites of cosmic ray acceleration up to a few PeV, though conclusive experimental evidence is still lacking.} {Hadron acceleration is expected to increase the shock compression ratio, providing higher postshock densities, but X-ray emission from shocked ambient medium has not firmly been detected yet in remnants where particle acceleration is at work. We exploited the deep observations of the XMM-Newton Large Program on SN 1006 to verify this prediction.} {We performed spatially resolved spectral analysis of a set of regions covering the southeastern rim of SN~1006. We studied the spatial distribution of the thermodynamic properties of the ambient medium and carefully verified the robustness of the result with respect to the analysis method.} {We detected the contribution of the shocked ambient medium. We also found that the postshock density of the interstellar medium significantly increases in regions where particle acceleration is efficient. Under the assumption of uniform preshock density, we found that the shock compression ratio reaches a value of $\\sim6$ in regions near the nonthermal limbs.} {Our results support the predictions of shock modification theory and indicate that effects of acceleration of cosmic ray hadrons on the postshock plasma can be observed in supernova remnants.} ",
        "introduction": "Supernova remnant (SNRs) are strong candidates for being the main source of energetic cosmic rays up to at least $3\\times10^{15}$ eV (\\citealt{be87,bv07}). X-ray synchrotron emission from high-energy electrons accelerated at the shock front up to TeV energies was first observed in SN~1006 \\citep{kpg95} and then in other young SNRs \\citep{rey08}, but a firm detection of high-energy hadrons is still lacking. %  Recently, TeV emission has been detected in SN~1006 \\citep{aaa10} and in a handful of SNRs showing bright nonthermal X-ray emission (\\citealt{aab07a,aab07,aaa07m,aaa09,aaa11}). The origin of the gamma-ray emission can be leptonic (i.~e., inverse Compton from the accelerated electrons) and$/$or hadronic (i.~e., proton-proton interactions with $\\pi^0$ production and subsequent decay). The hadronic scenario would directly prove that SNRs can accelerate cosmic rays up to PeV energies. Unfortunately, it is not easy to unambiguously ascertain the origin of the gamma-ray emission, as, for example, in RXJ1713.7-3946, where both hadronic \\citep{bv10} and leptonic \\citep{aaaa11,eps10,psr10} scenarios have been invoked to explain the observed multiwavelength observations. However, different indications concur in supporting the presence of high-energy hadrons in some young SNRs (e.~g., Tycho, \\citealt{ekh11,mc12}, and RCW 86, \\citealt{hvb09}). As for SN~1006, a pure leptonic model is consistent with the observations, but a mixed scenario that includes leptonic and hadronic components also provides a good fit to the gamma-ray data \\citep{aaa10}.  An alternative way to reveal hadron acceleration in SNRs is to probe its effects on the shock dynamics. The nonlinear back-reaction of high-energy particles on background plasma is predicted to strongly modify the shock properties by increasing the shock compression ratio and decreasing the postshock temperature with respect to the expected Rankine-Hugoniot values \\citep{be99,deb00,bla02,vyh10}. This effect is known as ``shock modification\". The observational test of these predictions requires accurate diagnostics of the thermal X-ray emission from the shocked interstellar medium (ISM). However, the ISM contribution in the X-ray spectra of remnants where particle acceleration is efficient is typically masked out by the bright synchrotron emission and by the thermal emission from shocked ejecta, as in RXJ1713.7-3946 \\citep{abd09},  Vela Jr \\citep{paf10}, G1.9+0.3 \\citep{brg10}, and Tycho \\citep{chb07}, and has not been firmly detected so far. In SN~1006 the quest for X-ray emission from shocked ISM is also ongoing \\citep{abd07,mbi09}.  SN~1006 exhibits a morphology characterized by two opposed radio, X-ray, (and $\\gamma-$ray) bright limbs dominated by nonthermal emission (bilateral SNR) and separated by an inner region of low surface brightness and soft, thermal X-ray emission. Thermal X-ray emission has been associated with shocked ejecta \\citep{abd07,mbi09}, consistent with the detection of Fe-rich plasma in the interior of the shell \\citep{ykk08}. A model with a nonthermal component plus a thermal component associated with the ejecta can fit the archive \\emph{XMM-Newton} spectra. An additional thermal ISM component is not needed from a statistical point of view \\citep{abd07,mbi09}, though some hints about ISM temperature and density ($kT_{\\rm ISM}\\sim 1.5-2$ keV, $n<0.2$ cm$^{-3}$) can be inferred \\citep{abd07}. \\emph{Suzaku} spectra of SN~1006 have been modeled with three thermal components only by assuming that the soft component originated in the ISM; nevertheless, it was not possible to exclude that the O lines, which dominate the soft emission, originate in the ejecta \\citep{ykk08}. A firm detection of X-ray emission from the ISM is therefore still lacking.  Indirect evidence for shock modification in SN~1006 has been obtained by measuring the distance, $D_{SFCD}$, between the shock front and the contact discontinuity, which is expected to be smaller in nonthermal limbs where particle acceleration is more efficient. $D_{SFCD}$ is instead almost the same all over the shell (even in regions dominated by thermal emission), though it is much smaller than that expected from unmodified shocks \\citep{chr08,mbi09}. Recently, 3-D magneto-hydrodynamic simulations have shown that this small distance can be naturally explained by ejecta clumping, without invoking shock modification \\citep{obm12}. Therefore, the small value of $D_{SFCD}$ is not a reliable indicator of hadronic acceleration. %  Here new, deep observations of the \\emph{XMM-Newton} SN~1006 Large Program allow us to present the detection of X-ray emission from shocked ISM in SN~1006 and to show indications for the effects of hadron acceleration on the postshock density of the ambient medium.  The paper is organized as follows: in Sect. \\ref{DP}, we describe the data analysis procedure and the background subtraction; in Sect. \\ref{Results}, we show the results of the spatially resolved spectral analysis; and, finally, we discuss our conclusions in Sect. \\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}  We have analyzed a set of deep \\emph{XMM-Newton} observations of SN~1006 that have allowed us to address important issues.  We have detected X-ray emission from the shocked ISM in the southeastern rim of SN~1006. It was not possible to explore the conditions of the shocked ISM in the nonthermal limbs because the higher contribution of the synchrotron radiation did not allow us to obtain tight constraints. Nevertheless, in regions $a-h$, the high statistics of the new observations showed that the addition of a thermal, underionized component improves the quality of our spectral modeling. A further indication that the association of this additional component with the ISM is sound is that its contribution to the nitrogen line emission is significant. Nitrogen line emission cannot be associated with ejecta in remnants of Type Ia supernovae and should originate in the shocked ISM. As shown in Fig. \\ref{fig:spec}, the contribution of the ejecta component dominates below 0.5 keV (i.~e., the band of the N lines). However, we found that the fit of our model to the spectra does not change significantly if we set the nitrogen abundance in the ejecta component to zero, while by putting N=0 in the ISM component, we obtained significant residuals around the nitrogen line. This indicates that in this band the contribution of the ISM component to the nitrogen emission lines is, as expected, stronger then the contribution of the ejecta component to the same line.  Figures \\ref{fig:dens} and \\ref{fig:dens03-2} clearly indicate that, under the assumption of constant preshock density, the shock compression ratio increases (up to $\\sim5.5-6$) in regions where nonthermal emission is stronger. This is qualitatively consistent with the expectations from modified shock models. In particular, such a high shock compression ratio indicates that the fractional cosmic-ray pressure in the postshock region is $\\sim30\\%$ (\\citealt{vyh10}) near the nonthermal limbs.  Though a detailed theoretical modeling of the expected azimuthal profile of the shock compression ratio in SN~1006 is not available in the literature, we compared the observed ISM density profile with that predicted by the shock modification model provided by \\citet{vbk03} and shown as dashed$/$dotted blue curves in Fig. \\ref{fig:dens}. The solid blue curve was instead derived by assuming that the injection efficiency\\footnote{The injection efficiency is the fraction of particles injected in the acceleration process.}, $\\eta$, is proportional to the radio flux divided by $B^{3/2}$ (and has its minimum at $\\eta=5\\times10^{5}$ in region $e$), and adopting the magnetic field model MF2 of \\citet{pbm09}. The green line shows the constant compression-ratio scenario, which corresponds to no shock modification. For all the curves, the relationship between the injection efficiency and the shock compression factor has been obtained by following \\citet{fdb10}. The shock modification models predict an azimuthal trend similar to the measured one, though they cannot fit all the points. In particular, the models shown in Fig. \\ref{fig:dens} slightly underpredict the observed density variations. However, we point out that a complete, self-consistent model of hadron acceleration at oblique shocks is not available yet and that current models do not provide tight theoretical constraints, though they allow us to get qualitative indications. Moreover, the models available do not include a dependence of the compression ratio on the maximum energy of the accelerated hadrons which may be higher near the nonthermal limbs. The cut-off frequency of the synchrotron emission from ultrarelativistic electrons is, in fact, much higher in the limbs than in the center and elsewhere in the rim \\citep{rbd04,mbi09}. This indicates that electrons are accelerated to higher energies in the limbs and suggests that protons may also be accelerated to higher energies.  In conclusion, we have detected the contribution of the shocked ISM to the X-ray emission in the southestern rim of SN~1006. We also found an azimuthal trend of the ISM postshock density that suggests that cosmic ray acceleration is at work at the shock front of SN~1006 and modifies the shock compression ratio. The results we have obtained in the southeastern rim of SN 1006 therefore indicate the presence of shock modification by particle acceleration. A definitive proof will require more data at higher spatial resolution. These data are necessary to extract spectra from narrow regions between the shock front and the contact discontinuity in order to better isolate the contribution of the shocked ISM and to confirm the effect that we have detected in this work."
    },
    "1208/1208.2893_arXiv.txt": {
        "abstract": "The mechanism causing breaks in the radial surface brightness distribution of spiral galaxies is not yet well known. Despite theoretical efforts, there is not a unique explanation for these features and the observational results are not conclusive. In an attempt to address this problem, we have selected a sample of 34 highly inclined spiral galaxies present both in the Sloan Digital Sky Survey and in the \\textit{Spitzer} Survey of Stellar Structure in Galaxies. We have measured the surface brightness profiles in the five Sloan optical bands and in the  3.6$\\mu m$ \\textit{Spitzer} band. We have also calculated the color and stellar surface mass density profiles using the available photometric information, finding two differentiated features: an innermost \\textit{break radius} at distances of $\\sim 8 \\pm 1$ kpc [$0.77 \\pm 0.06$ $R_{25}$] and a second characteristic radius, or \\textit{truncation radius}, close to the outermost optical extent ($\\sim 14 \\pm 2$ kpc [$1.09 \\pm 0.05$ $R_{25}$]) of the galaxy. We propose in this work that the breaks might be a phenomena related to a threshold in the star formation, while truncations are more likely a real drop in the stellar mass density of the disk associated with the maximum angular momentum of the stars. ",
        "introduction": "\\label{sec:intro}  The vast majority of spiral galaxies do not follow a radial surface brightness decline mimicking a perfect exponential law as proposed by \\citet{patt40,devau58} \\& \\citet{free70}. Depending on the shape of the surface brightness distribution, a classification for face-on galaxies has been developed by \\citet{erw05} and \\citet{pt06}. It distinguishes three different types of profiles. Type I (TI) is the classical case, with a single exponential describing the entire profile; Type II (TII) profiles have a downbending brightness beyond the break. Type III (TIII) profiles are characterized by an upbending brightness beyond the break radius. Relative frequencies for each type are 10\\%, 60\\% and 30\\% \\citep{pt06} in the case of late-type spirals.  Photometric studies of TII galaxies reveal that the radial scalelength of the surface brightness profiles changes when a characteristic radius ($\\sim$10~kpc) is reached. This so called break radius is described in several studies of face-on galaxies \\citep{erw05,pt06,erw08} and is also found if galaxies are observed in edge-on projections \\citep{van82,grijs01,kruit07}. Using faint magnitude stars, \\citet{fer07} found also a break in the surface brightness profile of M~33. On the contrary, no break was detected by star counting neither in NGC~300 \\citep{bh05,vla09} nor in NGC~7793 \\citep{vla11}. Works on galaxies beyond the nearby Universe \\citep{per04,tp05,azzo08} show the presence of a break at redshifts up to z~$\\sim$~1. These results suggest that breaks, once formed, must have been stable for the last 8 Gyrs of galaxy evolution. Cosmological simulations \\citep{gov07,mar09} support the idea of a break in the light distribution of disk galaxies as well.  Different mechanisms explaining the origin of TII breaks have been proposed. All those theories can be sorted into two families. A possible scenario is that the break could be located at a position where a threshold in the star formation occurs \\citep{fall80,kenn89,elme94,scha04,elme06}. A change in the stellar population would thus be expected at the break radius, but not necessarily a downbending of the surface mass density profile as shown by some simulations \\citep{pat09,mar09}. Supporting this, \\citet{bakos08} found that in a sample of 85 face-on galaxies (see Pohlen \\& Trujilo, 2006), the ($g'-r'$) Sloan Digital Sky Survey (SDSS) color profile of the TII galaxies is, in general, U-shaped, with a minimum at the break radius, hinting at a minimum also in the mean age of the stellar population at the break radius. This is in agreement with studies on resolved stellar populations across the break \\citep{jong07,rad12}. Furthermore, the surface mass density profile recovered from the photometry shows a much smoother behavior, in which the break is almost absent compared to the surface brightness profile. Numerical simulations \\citep{deba06,ros08,pat09,mar09} reveal how the secular redistribution of the angular momentum through stellar migration can drive the formation of a break. In the papers of \\citet{ros08} and \\citet{pat09}, a minimum in the age of the stellar population is found at the break radius, in agreement with the results of \\citet{bakos08}. It is also remarkable that the mean break radius in the simulations of \\citet{ros08} is very close (2.6 $h_\\mathrm{r}$) to the observational result (2.5$\\pm 0.6$ $h_\\mathrm{r}$) measured by \\citet{pt06}, where the values are given in radial scalelength units ($h_\\mathrm{r}$). However, \\citet{van82} showed that the Goldreich-Lynden-Bell criterion for stability of gas layers offered a poor prediction for the truncation radius in their sample of edge-on galaxies. Focusing on these high inclined galaxies, \\citet{kruit87,kruit88} proposed an alternative scenario where the break could be related to the maximum angular momentum of the protogalactic cloud, leading to a break radius of around four or five times the radial scalelength. This mechanism would also lead to a fast drop in the density of stars beyond the break. Observations of edge-on galaxies, such as those by \\citet{van82}, \\citet{bart94} or \\citet{kreg02}, place the break at a radius around four times the radial scalelength, supporting this second explanation.  There is a clear discrepancy between the positions of the breaks located in the face-on view compared to those breaks found in the edge-on perspective. In the face-on view, the breaks are located at closer radial distances of the center than in the edge-on cases. Are the two types of breaks the same phenomenon? On the one hand, the edge-on observations tend to support the idea of the maximum angular momentum as the main actor in the break formation. On the other hand, face-on galaxies, supported by numerical simulations, favor a scenario where the break is associated with some kind of threshold in the star formation along the disk. Only a few studies have attempted to give a more global vision of the problem \\citep{poh04,poh07} by deprojecting edge-on surface brightness profiles. \\citet{poh07} and \\citet{seb12} found that the classification exposed at the beginning of the current paper into TI, TII and TIII profiles is basically independent of the geometry of the problem. Consequently, the difference between the break radius obtained from the edge-on galaxies and the one from the face-on galaxies can not be related to the different inclination angle. More details on the current understanding of breaks and truncations can be found in the recent review by \\citet[$\\S$3.8]{van11}.  Despite truncations in edge-on galaxies and breaks in face-on galaxies have been traditionally considered equivalent features, our aim in the current paper is to understand the already noticed observational differences, proposing a global and self-consistent understanding of the breaks. We use images in six different filters (five SDSS bands and the S$^4$G 3.6$\\mu m$ band) to study the radial surface brightness distribution in a sample of 34 edge-on galaxies. We also measure the color and the stellar surface mass density profiles for each object, trying to constrain the most plausible mechanism for the break formation.  The layout of this work is as follows: in Section \\ref{sec:data} the characteristics of the sample are presented. Section \\ref{sec:prof} describes how the profiles are measured. The analysis is shown in Section \\ref{sec:ana} and then we discuss the results in Section \\ref{sec:discu}. The main conclusions listed in Section \\ref{sec:summ}. Tables referenced in the paper are in Appendix \\ref{sec:tables}.  Throughout, we adopt a standard $\\Lambda$CDM set of cosmological parameters ($H_0 = 70$ km  s$^{-1}$ Mpc$^{-1}$; $\\Omega_\\mathrm{M} = 0.30$;  $\\Omega_\\mathrm{\\Lambda} = 0.70$) to calculate the redshift dependent quantities. We have used AB magnitudes unless otherwise stated. ",
        "conclusions": "\\label{sec:summ} Using SDSS and S$^4$G imaging, we have found the following important aspects regarding the behaviour of surface brightness profiles in edge-on late-type spirals:  \\begin{enumerate}  \\item The majority of our galaxies (82 $\\pm$ 16\\%) show a TII surface brightness profile as those found in photometric studies of face-on galaxies, with the break occuring at a mean radial distance from the galactic center equal to 7.9 $\\pm$ 0.9 kpc.  \\item Truncations, previously described in edge-on galaxies as a quick drop in the surface brightness profile, have been found in 20 of the 34 galaxies in our sample. This drop, also observed in the stellar surface mass density profile, occurs at an average radial distance of 14 $\\pm$ 2 kpc.  \\item For many galaxies, breaks and truncations coexist as two differentiated features in the light distribution of the disks in spiral galaxies.  \\item Strong correlations exist between the truncation radius and the maximum rotational velocity, and the specific angular momentum of the disk. These correlations are, however, less strong in the case of breaks. This result reinforces the idea that breaks are more likely a phenomena related to a star formation threshold whereas truncations have a strong connection with the maximum angular momentum of the galaxy.  \\item Color and stellar surface mass density profiles are both very sensitive to the line of sight projection and to the presence of dust. Their interpretation in edge-on systems is not trivial and could be the cause of the differences between some of our results and those found in face-on works (e.g. correlations between $M_\\mathrm{B}$ and break radius).  \\end{enumerate}  \\small{ \\footnotesize{\\textit{Acknowledgments.}} We would like to thank Alexandre Vazdekis and Jes\\'us Falc\\'on-Barroso for their useful comments. We especially thank the referee, Prof. Piet van der Kruit, for providing very constructive and precise comments on this article.  This work has been supported by the Programa Nacional de Astronom\\'ia y Astrof\\'isica of the Spanish Ministry of Science and Innovation under grant AYA2010-21322-C03-02. We acknowledge financial support to the DAGAL network from the People Programme (Marie Curie Actions) of the European Union's Seventh Framework Programme FP7/2007-2013/ under REA grant agreement number PITN-GA-2011-289313.  Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy Office of Science. The SDSS-III web site is http://www.sdss3.org/. SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, University of Cambridge, Carnegie Mellon University, University of Florida, the French Participation Group, the German Participation Group, Harvard University, the Instituto de Astrof\\'isica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, Max Planck Institute for Extraterrestrial Physics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University.  This research has made use of NASA's Astrophysics Data System. We acknowledge the usage of the HyperLeda database and the  NASA/IPAC Extragalactic Database (NED), operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.  }"
    },
    "1208/1208.0896_arXiv.txt": {
        "abstract": "{} {We aim to investigate the relation between the long-term flux density and the position angle (PA) evolution of inner-jet in blazars.} {We have carried out the elliptic Gaussian model-fit to the `core' of 50 blazars from 15 GHz VLBA data, and analyzed the variability properties of three blazars from the model-fit results.} {Diverse correlations between the long-term peak flux density and the PA evolution of the major axis of the `core' have been found in $\\sim$ 20\\% of the 50 sources. Of them, three typical blazars have been analyzed, which also show quasi-periodic flux variations of a few years (T). The correlation between the peak flux density and the PA of inner-jet is positive for S5~0716+714, and negative for S4~1807+698. The two sources cannot be explained with the ballistic jet models, the non-ballistic models have been analyzed to explain the two sub-luminal blazars. A correlation between the peak flux density and the PA (with a T/4 time lag) of inner-jet is found in [HB89]~1823+568, this correlation can be explained with a ballistic precession jet model. All the explanations are based mainly on the geometric beaming effect; physical flux density variations from the jet base would be considered for more complicated situations in future, which could account for the no or less significance of the correlation between the peak flux density and the PA of inner-jet in the majority blazars of our sample.} ",
        "introduction": "Blazars are thought of the extreme of active galactic nuclei (AGNs), which consist of flat-spectrum radio quasars and BL Lac objects. They are variable in almost whole electromagnetic spectrum from radio to gamma-ray. The extreme properties of blazars are mainly attributed to their jets which are pointing closely to our line of sight, according to the unified model of AGN (Urry \\& Padovani 1995). With very long baseline interferometry (VLBI) -- the highest angular resolution to date in astronomy -- blazars can be resolved into a bright core and often one-sided jet components in parsec scale or sub-parsec scale (Lister et al. 2009a). In the VLBI scales, the blazar jets frequently show fast outward motions, with an apparent advancing speed of greater than the speed of light -- namely super-luminal motion (Rees 1966). This is due to the so-called relativistic Doppler boosting -- relativistic beaming effect (e.g. Cohen et al. 2007; Hovatta et al. 2009). Such high flux variability and VLBI structure changes together provide us an effective means to study the inner-jet property and probably the central engine's property of the blazars in further.  Recently, Britzen et al. (2009) found there seems to be evidence for an apparent stationarity of jet components (with regard to their core separation) with time, in blazar S5~0716+714, while the inner-jet components exhibit strong changes in their position angles. This indicates that the jet components move non-radially (or non-ballistic) with regard to their position angle. These authors attributed such sub-luminal motions to a geometric origin, although the geometric models were not well defined yet. Furthermore, they found that the long-term flux density and the position angle of the jet in 0716+714 show a significant positive correlation. The similar phenomena of the apparent stationarity of jet components with time have also been reported in BL Lac objects S5~1803+784 and PKS~0735+178, by Britzen et al. (2010a, 2010b). From CJF sample (Caltech-Jodrell Bank flat-spectrum sample of radio loud active galaxies), Karouzos et al. (2012) investigated the morphologies and pc-scale jet kinematics of more than 200 sources by using the so-called jet ridge-line method. They found that about a half of the sample show jet widths $>10\\degr$, with BL Lac jet ridge lines showing significantly larger apparent widths than both quasars and radio galaxies, indicating that the BL Lacs have more significant position angle changes of jets in their evolution.  From the study of AGN jet kinematics both in the 2 cm survey (Kellermann et al. 2004) and in the radio reference frame image database (RRFID), about one-third of well-measured component trajectories are non-radial (Piner et al. 2007). In the MOJAVE (monitoring of jets in active galactic nuclei with VLBA experiments) sample, 166 out of the 526 robust components (32\\%) show non-radial motions, indicating that non-radial motion is a common feature of the jet flow in blazars (Lister et al. 2009a, 2009b). Three of these non-radial components might be considered inward at the 3$\\sigma$ level (Lister et al. 2009b), this could be due to a projection effect of curved jets. With the acceleration measurements of the components of MOJAVE sources, Homan et al. (2009) found the perpendicular accelerations are closely linked with non-radial motions, and about half of the components show `non-radial' motions; but the parallel accelerations are generally larger than the perpendicular accelerations with respect to the component velocities.  To further investigate the relation, which was first discovered in 0716+714 by Britzen et al. (2009), between the long-term flux density and the position angle (PA) evolution of inner-jet in blazars, we model-fitted the `core' components of 50 core-dominated balzars from the MOJAVE data released up to 2011 (Lister et al. 2009a). We aimed at finding more sources which showing correlations between the long-term flux density and the PA evolution of inner-jet. We also analyzed the inner-jet models to explain the typical correlations, if any, in the model-fit results. ",
        "conclusions": "We summarize our findings from the three BL Lac objects. There exhibit diverse correlations between the peak flux density and the PA of the major axis of `core' component in the three sources. They also show quasi-periodic variations of the peak flux density, around a few years.  To explain the correlations and periodicity mentioned above, we have to find the models related to the inner-jet process in the radio loud AGNs. In general, from the literature there are several parameters/factors related to jet models, which should be clarified and clearly defined, e.g., `ballistic' (we define as a single jet component always moving along a straight line once launched from the jet nozzle), `non-ballistic' (we define as a single jet component moving in a curved way), `precession' (this has been confused in the literature, we define it as a precession `jet' in terms of a jet flow (by no means a single jet component), in which every single jet component is ballistically launched from a precession jet-nozzle), `non-precession' (we define as a jet nozzle has no precession). There are four meaningful hybrids of the parameters mentioned above, i.e., ballistic jets with a non-precession nozzle (B+nP), ballistic jets with a precession nozzle (B+P), non-ballistic jets with a non-precession nozzle (nB+nP), non-ballistic jets with a precession nozzle (nB+P). The quasi-periodic changing of the PA of inner-jet in the three sources, could be explained with the hybrid models except the first one (B+nP). A precession of the jet nozzle can be due to e.g., double super-massive black holes or other reasons (see Britzen et al. 2010a and references therein). For the non-ballistic jet, models have been proposed e.g., Gong et al. (2011), also see Lister et al. (2009b) and Britzen et al. (2009, 2010a, 2010b) for discussions. In these models, a non-ballistic motion of jet is largely related to the jet interaction with ambient matter. In the following, we will use the hybrid inner-jet models we classified, for the three sources.  Firstly, we try to apply these models to 1823+568, after analysis below, the B+P model (i.e. ballistic jets with a precession jet nozzle) is applicable to this source, as shown in the schematic diagram of Fig.~\\ref{fig6}. We are nearly face-on the jet with the estimated jet viewing angle of $\\sim$8.4\\degr (Savolainen et al. 2009), and the inner-jet PA is oscillating between $-$151\\degr and $-$165\\degr (see Table~\\ref{tab1} and Fig.~\\ref{fig4}). In Fig.~\\ref{fig6}, our line of sight is close to the phase `A' of a precession jet flow. At the phase `A', the relativistic beaming effect is the strongest, so the flux density peaks; on the contrast, the phase `C' of precession jet flow is far from our line of sight, so the flux density is the lowest; at the phase `B' or `D' the flux density has a middle value due to the beaming effect. For the periodic peak flux variations (T$\\sim$7.0 yr in Table~\\ref{tab1}), and the correlation between the peak flux density and the PA (with a T/4 time-lag) of inner-jet, the precession of the jet flow must be counter-clockwise, i.e. through the phases `B$\\rightarrow$A$\\rightarrow$D$\\rightarrow$C$\\rightarrow$B' in Fig.~\\ref{fig6}. Such that the PA peaks (a minus value) at the phase `B' while the flux density is at a middle level; and the PA has a middle value at the phase `A' while the flux density peaks; and the PA has the smallest (minus) value at the phase `D' while the flux density is at a middle level; and the PA has a middle value at the phase `C' while the flux density is the lowest (also see, Fig.~\\ref{fig4}). So the correlation between the peak flux density and the PA (with the T/4 time lag) of inner-jet can be well explained. This `B+P' model seems to support the super-luminal motions of 20.8~c (Lister et al. 2009b) or 9.4~c (Savolainen et al. 2009) in this source, because the ballistic motions have generally faster radial speed than that of the non-ballistic motions. The origin of the periodic precession of $\\sim$7.0 yr in the jet nozzle still needs to be investigated in future.  For the source either S5 0716+714 or 1807+698, the correlations between the peak flux density and the PA of inner-jet cannot be explained with a ballistic jet model. For 0716+714, the inner-jet PA is oscillating between 6.6\\degr and 43\\degr (see Table~\\ref{tab1} and Fig.~\\ref{fig2}). We are nearly face-on the jet with the estimated jet viewing angle of $\\sim$5.3\\degr (Savolainen et al. 2009). The significant positive correlation (Table~\\ref{tab1}) indicates that the flux density peaks while the PA peaks (i.e. far from North) simultaneously (see Fig.~\\ref{fig2}), so it is impossible that the jet at the phase `A' in the schematic diagram (Fig.~\\ref{fig7} left) has a middle-level of flux density according to the beaming effect within a ballistic model. Non-ballistic jet models have to be invoked, in which a single jet component is moving in a curved way. In the non-ballistic scenario, regardless the jet nozzle is in precession or not, the jet must have a non-radial motion. For 0716+714, in the schematic diagram Fig.~\\ref{fig7} (left), for a non-ballistic jet moving along the counter-clockwise helical trajectory `A$\\rightarrow$B$\\rightarrow$C$\\rightarrow$D', i.e. the curved way through which the jet at the phase `D' is pointing the closest to our line of sight, such that the flux density peaks while the PA peaks too. The peak flux density has a middle value while the PA has a middle value too at the phase `A', and so on, the positive correlation between the peak flux density and the PA at other phases `B', `C' can also be explained with the counter-clockwise non-ballistic helical jet model (see Fig.2 and Fig.~\\ref{fig7} left). The non-ballistic jet model will lead to a less outward motion speed than that of a ballistic jet model, assuming the originally-launched jet speed is the same. This is consisting with the sub-luminal motion or the `stationary scenario' of 0716+714, as proposed by Britzen et al. (2009). In general, a non-ballistic jet is likely to be formed, through the shocked jet plasma working on the surrounding medium (Gong et al. 2011), and a helical like trajectory of the jet ridge-line would be expected.  For 1807+698, the inner-jet PA is oscillating between $-$99\\degr and $-$110\\degr (see Table~\\ref{tab1} and Fig.~\\ref{fig3}). We are not well face-on the jet for the estimated jet viewing angle of $\\sim$45.3\\degr (Savolainen et al. 2009). A clockwise non-ballistic jet motion can explain the anti-correlation (i.e. a T/2 time lag) between the peak flux density and the PA of inner-jet, i.e. a jet is moving through `A$\\rightarrow$B$\\rightarrow$C$\\rightarrow$D' in Fig.~\\ref{fig7} right, the jet through at the `D' is pointing the closest to our line of sight. Such that the flux density peaks due to beaming effect at the `D' while the PA has a smallest (minus) value, and at the phases `A' and `C', both the peak flux density and the PA have middle values, and at the phase `B' the flux density is the lowest while the PA has the largest (minus) value. The source jet is sub-luminal (0.1~c, Lister et al. 2009b), which is the non-ballistic jet model expected. The jet viewing angle of 1807+698 is relatively large (45.3\\degr), so that the beaming effect is reduced. This may have led to the relatively low significance of the correlation between the peak flux density and the PA of inner-jet of this source (see Table~\\ref{tab1}), and the lower significance of the estimated period of the flux variations.  Among the three sources, the periodicity of flux variability and the correlation between the peak flux density and the T/4-laged PA of inner-jet in 1823+568 is a purely geometric effect according to the ballistic precession jet model `B+P'. The quasi-periodicity of flux variability and the correlations between the peak flux density and the PA of inner-jet in 0716+714 and 1807+698, as the non-ballistic models considered above, are also mostly caused by the geometric beaming effect. For the non-ballistic models in particular, however, even if the precession of the jet nozzle is invoked, the repeating periodic flux variability frequently found in blazars still needs to be fully explained. A kind of physical (non-geometric) outburst from the jet nozzle, or a physically periodic flux variations from the jet base would be considered for 0716+714 and 1807+698 in the future, which may be related to perturbations induced by the changes of the accretion mode or by a periodic intervening of a companying object. Multi-bands study of the three blazars on e.g., the SED, is also needed, like that recently being investigated for balzars statistically (see, Meyer et al. 2011; Lister et al. 2011)."
    },
    "1208/1208.3198_arXiv.txt": {
        "abstract": "We present a simple model for the relationship between quasars, galaxies, and dark matter halos from $0.5<z<6$.  In the model, black hole (BH) mass is linearly related to galaxy mass, and galaxies are connected to dark matter halos via empirically constrained relations.  A simple ``scattered'' light bulb model for quasars is adopted, wherein BHs shine at a fixed fraction of the Eddington luminosity during accretion episodes, and Eddington ratios are drawn from a lognormal distribution that is redshift-independent.  This model has two free, physically meaningful parameters at each redshift: the normalization of the $\\Mbh-\\Mgal$ relation and the quasar duty cycle; these parameters are fit to the observed quasar luminosity function (LF) over the interval $0.5<z<6$.  This simple model provides an excellent fit to the LF at all epochs, and also successfully predicts the observed projected two-point correlation of quasars from $0.5<z<2.5$.  It is significant that a {\\it single} quasar duty cycle at each redshift is capable of reproducing the extant observations.  The data are therefore consistent with a scenario wherein quasars are equally likely to exist in galaxies, and therefore dark matter halos, over a wide range in masses.  The knee in the quasar LF is a reflection of the knee in the stellar mass-halo mass relation.  Future constraints on the quasar LF and quasar clustering at high redshift will provide strong constraints on the model.  In the model, the autocorrelation function of quasars becomes a strong function of luminosity only at the very highest luminosities, and will be difficult to observe because such quasars are so rare.  Cross-correlation techniques may provide useful constraints on the bias of such rare objects.  The simplicity of the model allows for rapid generation of quasar mock catalogs from N-body simulations that match the observed luminosity function and clustering to high redshift. ",
        "introduction": "\\label{sec:introduction}  Quasars are among the most luminous astrophysical objects, and are believed to be powered by accretion onto supermassive black holes \\citep[e.g.][]{Sal64,Lyn69}.  They have become a key element in our current paradigm of galaxy evolution \\citep[e.g.,][]{Spr05, Croton06, Hop08}, and essentially all spheroidal systems at present harbor massive black holes \\citep{KorRic95}, the masses of which are correlated with many properties of their host systems.  Despite their importance, and intense theoretical activity, a full theory of the coevolution of galaxies and quasar eludes us.  The current paradigm assumes that every galaxy initially forms in a gas-rich, rotationally-supported system.  Once the dark matter halo grows to a critical scale some event, most likely a major merger \\citep{Car90, HaiLoe98, CatHaeRee99, KauHae00, Spr05, Hop06, Hop08} or instability in a cold-stream fed disk \\citep{DiM12}, triggers a period of rapid, obscured star formation, the generation of a stellar bulge and a growing black hole (BH).  Eventually the accreting BH becomes visible as a quasar, and soon after the star formation is quenched on a short timescale, perhaps via radiative or mechanical feedback from the BH \\citep[e.g.][]{Sil98, Kin03, WyiLoe03, Sha09, Nat12, AleHic12}. Understanding the details of this picture remains an active area of research.  Phenomenological models for quasar demographics often adopt power-law relations between quasars, galaxies, and dark matter halos \\citep[e.g.,][]{EfsRee88, Car90, WyiLoe02, WyiLoe03, HaiCioOst04, Mar06, Lid06, Croton09, She09, BooSch10}.  In these models, the duty cycle of quasars is tuned to match the observations, and a generic conclusion is that the duty cycle is a strong function of halo mass or quasar luminosity, peaking at a halo mass of $10^{12-13}\\Msun$. However, these previous models do not incorporate constraints provided by the galaxy stellar mass function over the interval $0<z<6$.  And yet, a variety of lines of evidence suggest that the relation between halos and galaxies is highly non-linear, with a characteristic peak in galaxy formation efficiency at a halo mass of $\\sim10^{12}\\Msun$ \\citep{Van03, Val04, Man06, Con09, Mos10, TruKlyPri11, BehWecCon12}. The aim of this paper is to incorporate empirically constrained relations between galaxies and halos into a simple model for quasar demographics.  We will demonstrate that a model constructed to match the observed galaxy stellar mass function implies a quasar duty cycle that is independent of galaxy and halo mass at each redshift.  This has important implications for physical models aimed at understanding the triggering of quasars and their connection to the evolution of galaxies.  The outline of the paper is as follows.  In \\S\\ref{sec:model} we describe the model, in \\S\\ref{sec:data} the model is compared to data, and a discussion is presented in \\S\\ref{sec:discussion}.  We conclude in \\S\\ref{sec:conclusions}.  Where necessary we adopt a $\\Lambda$CDM cosmological model with $\\Omega_m=0.28$, $\\Omega_\\Lambda=0.72$ and $\\sigma_8=0.8$.  Unless the $h$ dependence is explicitly specified or parametrized, we assume $h=0.7$.  Dark matter halo masses are quoted as $M_{\\rm vir}$ \\citep{Bryan98}.  Luminosities are quoted in Watts and magnitudes in the AB system, and stellar masses assume a \\citet{Cha03} stellar initial mass function. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Implications} \\label{sec:imp}  The success of our model in reproducing the basic demographics of quasars allows us to consider several implications that follow naturally within our framework.  In Figure \\ref{fig:mbh_mhalo} we show the best-fit model $M_h-\\Mbh$ relations from $z=0.5$ to $z=3.75$ (the relations above $z=3.75$ are highly under-constrained and so are not plotted).  As discussed above, the quasar LF places very weak constraints on the model relations at $\\log(M_h/\\Msun)>13.5$, and so one should interpret the model relations in Figure \\ref{fig:mbh_mhalo} with this in mind.  It is also worth pointing out that while the model formally allows for the existence of extremely massive BHs with $\\Mbh>10^{10}\\Msun$ residing within moderately massive halos, at high redshift such halos are very rare.  For example, at $z=4.75$ one expects only of order one halo with log$(M_h/\\Msun)>$13 per $10^9$ Mpc$^3$.  With average mass accretion histories for halos, we can evolve halos and hence their black holes through the relations shown in Figure \\ref{fig:mbh_mhalo}.  To do this we employ mass accretion histories presented in \\citet{BehWecCon12}, which provide excellent fits to the results of $N-$body simulations.  The resulting evolution in BH mass is shown in Figure \\ref{fig:mbh_growth} for three representative halo masses, and for both model choices for the evolution in the Eddington ratio.  In the model lower mass black holes are growing to lower redshift faster than higher mass black holes (this is sometimes referred to as BH downsizing). In the model with a constant $\\eta$, the BHs in the most massive halos lose mass below $z\\approx 1.5$, while in the varying $\\eta$ model all BHs grow, if only modestly, at all epochs.  This suggests that a model with evolving Eddington ratios may be necessary to ensure self-consistent evolution.  Models that enforce self-consistent growth of BHs should shed further light on this problem \\citep[e.g.,][]{Mer04, MerHei08, Sha09}.  \\begin{figure}[!t] \\begin{center} \\resizebox{3.5in}{!}{\\includegraphics{f10.eps}} \\end{center} \\caption{BH growth in the best-fit model from $z=3.75$ to $z=0.5$. Results are shown for two choices for the evolution in $\\eta$ (see the lower panel of Figure \\ref{fig:dutycycle}).  Notice that the constant $\\eta$ model produces massive BHs that lose mass at $z<1.5$, suggesting that one or more of the assumptions of this model are breaking down at low redshift.  In contrast, the varying $\\eta$ model produces realistic BH growth at all epochs.  In both models lower mass BHs grow more at late times compared to higher mass BHs, a phenomenon sometimes referred to as BH downsizing.} \\label{fig:mbh_growth} \\end{figure}   Figure \\ref{fig:mag_mhalo} shows the evolution of the halo mass for quasars of fixed luminosity.  The trend of lower $M_h$ at higher $z$ was already apparent in Figure \\ref{fig:mbh_mhalo}.  Figure \\ref{fig:mag_mhalo} also emphasizes how the range of halo masses for a fixed luminosity range narrows towards higher $z$.  This effect is in the opposite sense to models which tie the luminosity of quasars directly to halo properties \\citep[e.g.][]{Croton09}.  Our model is able to reproduce the observed $L$-independent clustering at low $z$ because the run of bias with halo mass also becomes shallower at low $z$ for the halo masses of interest.  The evolution of the LF shown in Figure \\ref{fig:lf} is driven by evolution in the $\\Mbh-\\Mgal$ and $\\Mgal-M_h$ relations and the evolution of the halo mass function (evolution in the $L_Q-\\Mbh$ relation is governed by evolution in $\\eta$).  The break in the model quasar LF arises primarily due to the shape of the $M_{\\rm gal}-M_h$ relation, and thus $L_\\star$ quasars live in halos near the peak of that relation, $M_h\\sim 10^{12}M_\\odot$.  The peak of the $\\Mgal-M_h$ relation changes very little with redshift \\citep[e.g.,]{BehWecCon12}, so that at fixed $M_h$ there is little change of $\\Mgal$ with $z$. However the luminosity of the break can evolve due to a combination of evolution in the $\\Mbh-\\Mgal$ relation or the Eddington ratio.  In our fiducial model $\\eta$ is constant and $\\Mbh\\propto (1+z)^2$ at fixed $M_{\\rm gal}$ and so the break in the luminosity function scales as $(1+z)^2$.  The faint-end slope of the model LF does not vary significantly, in good agreement with the data, and the overall normalization changes only modestly.  The major departure from pure luminosity evolution is the change in the slope of the bright end. The bright-end slope appears shallower at higher $z$ both because the data are probing closer to the (brighter) break of the LF and because the $\\Mbh-M_h$ relation becomes steeper at higher mass and redshift. We also note that the bright end of the model LF is strongly suppressed at $z<1.5$, and it is this suppression that is responsible for much of the drop in the quasar number density to lower redshift. The drop is a consequence of evolving Eddington ratios and the shallowing of the $\\Mbh-M_h$ relation at high mass, which is in turn driven by the very slow growth of massive galaxies at low redshift.  In fact, the model naturally reproduces the global rise and fall of the quasar number density over the interval $0.5<z<4.75$.  This follows simply from the evolution in the $\\Mgal-M_h$ and $L_Q-\\Mgal$ relations and the halo mass function; it does not require strong evolution in $t_Q$ at low $z$.  Specifically we do not invoke a decline in the cold gas fraction nor a decline in the major merger rate at $z<2$ in order to reproduce the observed decline in the abundance of quasars.  While these physical processes may ultimately be responsible for shaping the evolving relations between $L_Q$, $\\Mbh$, $\\Mgal$ and $M_h$, they do not appear explicitly in the model.   \\begin{figure}[!t] \\begin{center} \\resizebox{3.5in}{!}{\\includegraphics{f11.eps}} \\end{center} \\caption{Relation between halo mass and redshift for quasars of a fixed luminosity.  At low redshift the range of halo masses hosting quasars is very broad, but the distribution narrows substantially at high redshift.  This is simply a recasting of the relations shown in Figure \\ref{fig:mbh_mhalo}.} \\label{fig:mag_mhalo} \\end{figure}  Our model favors a different picture of how quasars inhabit massive halos compared to previous work.  Rather than having a preferred halo mass scale (around $10^{12}\\,M_\\odot$) for quasar activity, the present model allows for actively accreting black holes in a broad range of galaxy and halo masses.  The apparent preference for quasars to live in halos of $10^{12}\\,M_\\odot$ arises from the shape of the $\\Mgal-M_h$ relation, which reflects the well known fact that galaxy formation is most efficient in halos near $10^{12}M_\\odot$, along with the shape of the halo mass function.  Specifically, above the knee in the $\\Mgal-M_h$ relation halos become exponentially rare, while below the knee a large range in $\\Mgal$ maps into a small range in $M_h$. Thus, the {\\it average} halo mass of quasars will be close to the knee, despite the fact that quasars occupy a broad distribution of halo masses.  Due to its simplicity the model predicts the clustering of any population of quasars once the model parameters are fixed (e.g., by the observed LF).  Variation in the $L_Q-M_h$ scatter or $\\Mbh-M_{\\rm gal}$ slope do not strongly affect the predicted clustering, meaning that our model makes an essentially parameter-free prediction of the clustering of quasars as a function of luminosity and redshift. Overall the agreement between the predicted clustering and the observations is good, though there is a tendency for the model to slightly underpredict the observations and there is some tension at the highest redshifts.  This tension has been noted before -- the very high amplitude of clustering measured at $z\\sim 4$, in combination with the abundance, requires quasars to have a duty cycle approaching unity and almost no scatter in $L_Q$ at fixed $M_h$ \\citep{WhiMarCoh08, ShaWeiShe10}.  This is at odds with the very low number densities but power-law decline seen in the luminosity function at high $z$.  If the clustering measurements can be strengthened, possibly by cross-correlation of existing spectroscopic quasar samples with deeper photometric quasar or galaxy samples, then it will indicate that one of our assumptions is breaking down as we approach the era of rapid black hole growth in the early Universe.  We make no assumption about what triggers quasar activity, whether it be a major merger of two gas rich galaxies, a secular instability in a disk, or a critical halo mass.  In general it is quite difficult to translate abundance and clustering measurements into constraints on the underlying mechanisms that trigger quasar activity.  We can gain some insight by the fact that our duty cycle, or quasar lifetime, is relatively independent of redshift with a tendency to fall towards higher redshifts rather than rise.  If quasars are visible for a fixed, but short, time and are triggered by mergers then we expect $t_Q$ to scale with the merger rate \\citep[c.f.][]{Car90}.  The merger rate for halos, per halo, per unit redshift is relatively flat \\citep{LacCol93, Per03, CohWhi05, WetCohWhi09, FakMa09, Hop10b}, so if we can naively translate halo mergers into galaxy mergers we expect a rate (per unit time) scaling as $(1+z)H(z)\\propto (1+z)^{5/2}$ for $z\\gg 1$.  If quasars are visible for a constant interval after each merger then $t_Q\\propto 1+z$, which is not in good agreement with our best-fit relation.  Of course, galaxy merger rates can differ from halo merger rates.  A recent analysis by \\citet{Hop10a} suggests a rate per unit time scaling as $(1+z)^{1.5-2.0}$, which would lead to slower evolution in $t_Q$, as we observe.  Such agreement is not conclusive however, and we cannot rule out secular processes or a time-varying combination of multiple triggers.  \\subsection{Comparison to Previous Work}  The success of our model in explaining the basic demographics of quasars with relatively few, smoothly varying inputs goes a long way to explaining the manner in which forward modeling of the quasar population can succeed with relatively little fine tuning.  Both semi-analytic models \\citep[e.g.,][]{CatHaeRee99, KauHae00, KauHae02, VolHaaMad03, BroSomFab04, Granato04, Croton06, MonFonTaf07, Mal07, Bon09, Fan12, Hir12} and hydrodynamic simulations \\citep[e.g.,][]{Sij07, DeGraf11} adjust their subgrid models to ensure a reasonable match to the $\\Mgal-M_h$ relation over a broad redshift range, thus ensuring that galaxies populate halos in approximately the correct manner.  All of the models introduce a $\\Mbh-\\Mgal$ relation through either or a combination of common feeding mechanisms and feedback-limited BH growth.  As we have shown, with these two ingredients even simple lightcurve models are sufficient to match the basic demographics of quasars over a broad range of luminosity and redshift.  A good match to the data can be found for a wide range of scatter in $\\Mbh-\\Mgal$, or evolution in the scatter.  Conversely, if a model has difficulties reproducing the stellar mass function and its evolution then it will need to incorporate mass-dependent quasar physics that counteracts this deficiency in order to match the observed quasar properties.  By contrast, models that tie black hole properties directly to the underlying halo population need to introduce more complexity in order to reproduce the observed properties of quasars.  Recent examples include \\citet{Lid06}, \\citet{Croton09}, and \\citet{She09}, who all need to include mass- and redshift-dependent duty cycles to explain the shape and evolution of the quasar luminosity function.  While our model and theirs can produce qualitatively similar fits to the basic data, the explanations for the observed behaviors differ.  One of the most basic differences is the range of halos that host active quasars, and its evolution (discussed above).  This in turn affects how each model explains the evolution of the quasar LF and the luminosity-independence of quasar clustering.  Conventional wisdom is that the quasar duty cycle is required by the data to be a (strong) function of luminosity \\citep[e.g.][]{Ade05, Hop05, Lid06, Croton09, She09}.  In our model this is not the case. There are two major reasons for this.  The first is that we obtain a flattening of the $b(L)$ relation from the steepness of the $L_Q-M_h$ relation at low $L_Q$ and the second is the intrinsic scatter \\footnote{This scatter may arise due to time-dependent processes, i.e.~a high $L_Q$ object at the time of observation is not required to have always been or continue to be high $L_Q$.}  in that relation. Thus our model is not a ``light bulb'' model in the sense of \\citet{Hop05} and \\citet{Lid06}, who reserve that term for a model in which there is no scatter in $L_Q-M_h$.  However scatter in the $L_Q-M_h$ relation is {\\it expected\\/}, due to the observed scatter in $\\Mbh-\\Mgal$ and variation in Eddington ratios if from no other source; for this reason we refer to our model as a ``scattered'' light bulb model.  This expected level of scatter is enough to make $b(L)$ flat until extremely high $L$ or correspondingly low $\\bar{n}_Q$ \\citep[a similar behavior is seen in the model of][which is also not strictly a light bulb model in the above sense]{Croton09}.  For this reason we are able to obtain a model in which both the quasar lifetime and the quasar clustering are independent of $L$.  \\citet{Aird12} studied $X-$ray selected active galactic nuclei (AGN) as a function of galaxy mass at $z\\sim0.6$ and found no preference for AGN to be found in galaxies of a particular mass at fixed Eddington ratio, even for ratios as high as $\\eta\\gtrsim0.1$.  Their results suggest a duty cycle that does not depend strongly on galaxy mass, in excellent agreement with our results.  Finally, the apparent preference for quasars to live in halos of $10^{12}M_\\odot$, which has been noted by many authors, arises in our model from the shape of the $\\Mgal-M_h$ relation, which reflects the well-known fact that galaxy formation is most efficient in halos of $10^{12}M_\\odot$, in combination with the halo mass function.  Within the context of our model this cannot be taken as evidence for a merger driven origin to quasar activity, despite the fact that it is close to the small group scale where mergers may be more efficient, because it is not believed that the knee of the $\\Mgal-M_h$ relation is related to mergers.  \\subsection{Mock Catalogs}  While our intent has been to understand the quasar phenomenon, the model can also be used for the creation of mock catalogs from N-body simulations.  The simplicity of the model makes it easy to rapidly generate redshift-dependent quasar populations that have the correct luminosity function and clustering, given halo catalogs at the redshifts of interest.  The steps for creating such a catalog are straightforward:  \\begin{itemize}  \\item[1.] Adopt the redshift-dependent $\\Mgal-M_h$ relation from \\citet{BehWecCon12}, including scatter in $\\Mgal$ at fixed $M_h$.  \\item[2.] Use the $\\Mgal-\\Mbh$ relation from Equation \\ref{eqn:mbh_mgal} to assign BHs to galaxies, including 0.3 dex of scatter in $\\Mbh$ at fixed $\\Mgal$.  Fix the normalization of this relation to the local value, with no redshift evolution (because we advocate using the varying $\\eta$ model; see below).  \\item[3.] Randomly turn a fraction, $f_{\\rm on}$, of the BHs into active quasars.  As evident from Figure \\ref{fig:dutycycle}, the quasar lifetime is approximately constant at $3\\times10^7$ yr at $z<3$; we therefore advocate fixing $t_Q$ to this value.  One then determines the duty cycle via $f_{\\rm on}(z)=t_Q/t_H(z)$.  \\item[4.] For the active BHs, convert $\\Mbh$ into $L_Q$ using Equation \\ref{eqn:Lq_Mbh}, with an additional 0.3 dex of scatter in $L_Q$ at fixed $\\Mbh$.  Use the redshift-dependent Eddington ratio, $\\eta$, shown in the bottom panel of Figure \\ref{fig:dutycycle}.  We advocate using the varying $\\eta$ model because this model produces self-consistent BH growth at all redshifts (see Figure \\ref{fig:mbh_mhalo}).  \\end{itemize}  When simulations are populated with quasars in this way, the mock quasar LF and clustering will agree with all existing LF and clustering data at $z<3$.  In order to produce mock catalogs at higher redshifts one will need to include a drop in $t_Q$ as shown in Figure \\ref{fig:dutycycle}.  Such mock catalogs should prove useful in the context of ongoing and future planned surveys such as BOSS, bigBOSS, DES, Pan-STARRS, SUMIRE and LSST."
    },
    "1208/1208.3960_arXiv.txt": {
        "abstract": "{ We evaluate our capability to constrain the abundance of primordial tensor perturbations (primordial gravitational waves, PGWs) in cosmologies with generalized expansion histories in the epoch of cosmic acceleration. Forthcoming satellite and sub-orbital experiments probing polarization in the Cosmic Microwave Background (CMB) are expected to measure the $B-$mode power in CMB polarization, coming from PGWs on the degree scale, as well as gravitational lensing on arcminute scales; the latter is the main competitor for  the  measurement of PGWs, and is directly affected by the underlying expansion history, determined by the presence of a Dark Energy (DE) component. In particular, we consider early DE possible scenarios, in which the expansion history is substantially modified at the epoch in which the CMB lensing is most relevant.  We show that {the introduction of a parametrized DE } may induce a variation as large as  $30\\%$ in the ratio of the power of lensing and PGWs on the degree scale. We find that adopting the nominal specifications of upcoming satellite measurements, the constraining power on PGWs is weakened by the inclusion of the extra degrees of freedom, resulting in a reduction of about $10\\%$ of the upper limits on $r$ in fiducial models with no GWs, as well as a comparable increase in the error bars in models with non-zero tensor power. Moreover, we find that the inclusion of sub-orbital CMB experiments, capable of mapping the $B-$mode power up to the angular scales which are affected by lensing, has the effect of restoring the forecasted performances with a fixed cosmological expansion history corresponding to a cosmological constant. Finally, we show how the combination of CMB data with Type Ia SuperNovae (SNe), Baryonic Acoustic Oscillations (BAO) and Hubble constant allows to constrain simultaneously the primordial tensor power and the DE quantities in the parametrization we consider, consisting of present abundance and first redshift derivative of the energy density. { We compare this study with results obtained using the forecasted lensing potential measurement precision from CMB satellite observations, finding consistent results.}} ",
        "introduction": "\\label{sec:introduction}  Anisotropies in the Cosmic Microwave Background (CMB) represent one of the pillars of modern cosmology. Their statistical distribution, characterized primarily by the angular power spectrum, is consistent with a flat Friedmann Robertson Walker metric, expanding at a Hubble rate corresponding to about 70 km/s/Mpc, and composed by three main cosmological components, namely baryons and leptons representing about $4\\%$ of the total energy density, dark matter (DM, about $21\\%$) constituting the large part of the gravitational potential around collapsed or forming cosmological structures, and about $75\\%$ of a Dark Energy (DE) component, similar or coincident with a Cosmological Constant (CC), responsible for a late time phase of accelerated expansion. The primordial spectrum of density perturbations is almost scale invariant, corresponding to a Harrison-Zel'dovich power law shape in wavenumbers. Three satellites have been observing CMB anisotropies, the Cosmic Background Explorer \\citep{smoot_etal_1992}, the Wilkinson Microwave Anisotropy Probe \\citep{komatsu_etal_2011}, and Planck, which is expected to release cosmological data in early 2013 \\citep{the_planck_collaboration_2011}. Space observations will provide an all sky measurement of total intensity and polarization anisotropies down to a resolution of a few arcminutes, and a sensitivity of a few $\\mu$K per resolution element.  A number of sub-orbital experiments are planned and have been observing selected regions of the sky and frequency spectrum, looking for arcminute and sub-arcminute scale anisotropies in total intensity ($T$), as well as polarization\\footnote{see NASA ADS for the list of operating or planned sub-orbital CMB experiments.}. These observations will target most important and yet still undetected effects, dominating the curl component ($B-$modes) of the linear polarization pattern in CMB anisotropies \\citep{kamionkowski_etal_1997,zaldarriaga_seljak_1997}. On arcminute angular scales, the latter are dominated by the gravitational lensing of the anisotropies at last scattering by means of forming cosmological structures along the line of sight. A fraction of the gradient component of polarization ($E-$modes), dominating because powered by density fluctuations responsible for sub-degree acoustic oscillations at last scattering, is converted into $B-$modes by means of gravitational lensing \\cite{zaldarriaga_seljak_1998}. The power spectrum of the underlying DM distribution, and the primordial $E-$modes, produce a characteristic and broad lensing peak centered at $l\\simeq 1000$ in the $B-$mode power spectrum. Gravitational lensing has been recently detected in the damping tail of $T$ anisotropies by several groups \\cite{keisler_etal_2011,hlozek_etal_2012}, also cross-correlating the lensing with observed structures \\citep{sherwin_das_2012}, while  $B-$modes have not yet been detected, see \\citet{the_quiet_collaboration_2011} for the current upper limits. On the degree angular scales on the other hand, a primordial spectrum of tensor anisotropies or cosmological Gravitational Waves (GWs) would produce a narrow peak, rapidly vanishing on sub-degree angular scales, not supported by radiation pressure from massive particles, as is instead the case for $T$ and $E-$modes. On large angular scales, corresponding to several degrees in the sky, the decay of the GWs tail in the $B-$modes can be re-amplified though re-scattering onto electrons in the epoch of cosmic reionization. As for the case of lensing, only upper limits exist for the amplitude of PGWs through direct measurement of $B-$modes. The two effects compete for detection, and their different origin, primordial and linear for GWs, late and second order for lensing, has been exploited for designing separation techniques \\cite{hirata_seljak_2004}. {Furthermore, it has been analysed in the past how an accelerated expansion modifies the shape of the spectrum of PGWs as a result of propagation in a different space-time \\cite{zhangetal}. }  The lensing peak of $B-$mode anisotropies strongly depends on the history of cosmic expansion. It has been shown \\cite{acquaviva_baccigalupi_2006} that its amplitude may undergo variations of order $10\\%$ if the DE is dynamical at the epoch corresponding to the onset of acceleration, i.e. about $z\\simeq 1$, in which its actual amplitude is poorly constrained by existing measurements of the CMB or large scale structures. The $B-$mode lensing peak as a DE probe has been investigated by several authors \\citep{acquaviva_baccigalupi_2006,hu_etal_2006}, who in particular have shown how the lensing is capable of breaking the projection degeneracy affecting CMB anisotropies at the linear level, as it was recently confirmed in the context of lensing detection for sub-orbital $T-$mode experiments \\citep{sherwin_etal_2011}. On the other hand, the detection thresholds for cosmological GWs as well as the accuracy on DE constraints from CMB observations have  never been given by taking into account the full set of degrees of freedom, represented not only by the amplitude of primordial GWs, but also by those related to the expansion history,  parametrized through suitable DE models. The release of the latter degrees of freedom in the context of experiments aiming at the detection and characterization of $B-$mode anisotropies is expected to have a direct impact in the quoted detection thresholds of primordial GWs.  In this work we explore this issue, by investigating the sensitivity of forthcoming $B-$mode probes on primordial GWs abundance as well as DE dynamics when all the physical degrees of freedom shaping the $B-$mode power spectrum are considered and treated jointly. In this context, we consider in particular the interplay { between satellite measurements, accessing large scale polarization and extracting lensing mainly from $T$ and $E$ measurements, and the case of sub-orbital ones, directly probing lensing $B-$modes.} We will take as reference two among the most important forthcoming $B-$mode probes, EBEX \\citep{reichborn_kjennerud_etal_2011} and PolarBear \\cite{polarbear} as well as the all sky measurements featuring the nominal capabilities from Planck \\cite{:2006uk}.  This work is organized as follows. In Section \\ref{sec:spectrum} we describe the impact of a modified expansion history on the CMB lensing power. In Section \\ref{sec:analysis} we describe our set of simulated data as well as the reference experiments we consider. In Section \\ref{sec:results} we show and discuss our results, while in Section \\ref{sec:discussion} we draw our conclusions. ",
        "conclusions": "\\label{sec:discussion}  The primordial Gravitational Waves (PGWs) and lensing power constitute the dominant effects for the $B-$mode polarization in the anisotropies of the Cosmic Microwave Background (CMB). While the former is dominated by the physics of the early Universe, parametrized through the primordial tensor-to-scalar ratio $r$, the latter is instead due to structure formation, and thus influenced by the expansion rate at the epoch of the onset of cosmic acceleration. This, in turn, is dependent on the underlying  dynamics of the Dark Energy (DE). Despite both signals being present in the CMB $B-$modes, their joint measurement in terms of parameter estimation was never considered, and this work represents a first step in this direction.  { We first address the lensing relevance for constraining our parametrization of the expansion history, assuming no PGWs. We find comparable results when the lensing is extracted from $T$ and $E$ data and when the lensing is traced directly through lensing $B-$modes, by forthcoming satellite and sub-orbital data, respectively, {both for a Planck experiment and for a CMBpol experiment}. Focusing on the latter case, where the two processes directly compete for detection in $B-$modes, we quantify} the constraining power on the abundance of PGWs which is expected from combined forthcoming satellite and sub-orbital experiments probing CMB polarization in cosmologies with generalized expansion histories, parametrized through the present and first redshift derivative of the DE equation of state, $w_{0}$ and $w_{a}$, respectively. We find that in the case of pure satellite measurements, corresponding to the Planck nominal performance, the constraining power on GWs power is weakened by the inclusion of the extra degrees of freedom, resulting in an increase of about $10\\%$ of the upper limits on $r$ in fiducial models with no GWs, as well as a comparable increase in the error bars in models with non-zero tensor power. The inclusion of sub-orbital CMB experiments, capable of mapping the $B-$mode power up to the angular scales which are affected by lensing, has the effect of making such loss of constraining power vanishing below a detectable level. We interpret these results as a joint effect of the CMB and external datasets:  the former are able, in particular with the data from sub-orbital probes, to access the region of $B-$modes which is lensing dominated, and therefore sensitive to the DE abundance at the onset of acceleration; the latter, as the case of Type Ia SNe and the Hubble Space Telescope, are on the other hand strongly constraining the dynamics of cosmic expansion at present. By inspecting the constraints on all cosmological parameters, including those parametrizing the expansion history, we also show that the datasets we consider do not highlight new degeneracies in the parametrization we consider.  Our results indicate that the combination of satellite and sub-orbital CMB data, with the available external data useful to inquire the late time expansion history, can be used for constraining jointly the dynamics of the DE as well as the primordial tensor-to-scalar ratio, with no new degeneracies or significant loss of sensitivity in particular on $r$ with respect to the case in which a pure Cosmological Constant determines the late time cosmological expansion. Our assumptions of course include the nominal performance of these experiments, and no realistic data analysis consisting in the inclusion of foregrounds in the CMB data, as well as systematic errors. It would be interesting to further investigate this phenomenology in specific DE models, and considering the role of future surveys in giving more accurate constraints. \\newpage"
    },
    "1208/1208.4637_arXiv.txt": {
        "abstract": "Geostationary satellites are unique among orbital spacecraft in that they experience no appreciable atmospheric drag. After concluding their respective missions, geostationary spacecraft remain in orbit virtually in perpetuity. As such, they represent some of human civilization's longest lasting artifacts.  With this in mind, the EchoStar XVI satellite, to be launched in fall 2012, will play host to a time capsule intended as a message for the deep future. Inspired in part by the Pioneer Plaque and Voyager Golden Records, the EchoStar XVI Artifact is a pair of gold-plated aluminum jackets housing a small silicon disc containing one hundred photographs. The Cover Etching, the subject of this paper, is etched onto one of the two jackets. It is a temporal map consisting of a star chart, pulsar timings, and other information describing the epoch from which EchoStar XVI came. The pulsar sample consists of 13 rapidly rotating objects, 5 of which are especially stable, having spin periods $< 10$ ms and extremely small spindown rates.  In this paper, we discuss our approach to the time map etched onto the cover and the scientific data shown on it; and we speculate on the uses that future scientists may have for its data.  The other portions of the EchoStar XVI Artifact will be discussed elsewhere. ",
        "introduction": "In the early 1970s, a group of scientists and artists led by Carl Sagan and Frank Drake designed several artifacts and messages intended for extraterrestrial audiences. The first was the Pioneer Plaque placed on the {\\it{Pioneer 10}} and {\\it{11}} spacecraft, launched in 1972 and 1973, respectively (Sagan et al 1972).  In 1974, Frank Drake composed the ``Arecibo Message,'' a 1679-bit digital transmission beamed toward star cluster M13 to celebrate the renovation of the Arecibo Observatory. More detailed messages, including images and sounds,  were encoded onto the phonographic \ufffdGolden Records\ufffd and placed on the {\\it{Voyager 1}} and {\\it{2}} spacecraft, launched in 1977 toward deep space (Sagan et al 1978).  While geostationary spacecraft do not leave the solar system or even the Earth's environs, they inhabit time in a way that is similar to deep space probes. Located in the so-called ``Clarke Belt,'' some 36,000 km above the equator, geostationary communications satellites experience no appreciable atmospheric drag. These spacecraft have one of two fates. International regulations stipulate that satellite operators conduct end-of-life maneuvers to place derelict satellites into supersynchronous ``graveyard'' orbits (Collis 2009). In practice, however, operators regularly leave their spent satellites in loosely geosynchronous orbits. In both cases, these spacecraft remain orbiting virtually in  perpetuity (Walker \\& King-Hele 1986; Flohrer et al 2011).  As such, geostationary satellites are undoubtedly some of contemporary civilization's longest-lasting objects. As a scientific and cultural exercise, we can imagine scientists in the distant future looking in this zone for relics from past civilizations (Freitas \\& Valdes 1983, 1985). The EchoStar XVI Artifact represents an effort to communicate with these future scientists.  The full Artifact is composed of two interlocking gold-plated aluminum jackets housing a silicon disc called ``The Last Pictures,'' on which 100 photographs are nano-etched. The Artifact Cover Etching is located on the exterior of one of the two jackets.  In this paper, we focus on the Artifact Cover Etching, wherein we endeavored to create a map explaining {\\it{when}} the spacecraft originated, to whomever may find it in the future.  Specifically, we detail the several scientific and mathematical messages we developed for the Artifact Cover. The  photographic  contents of the Artifact's silicon disc are detailed elsewhere (Paglen 2012). ",
        "conclusions": "A comparison of the messages and the technologies on the current Artifact with those on the Pioneer and Voyager spacecraft indicates the rapid pace of technological advance over even just a few decades.  While we have endeavored to use putatively universal symbols, the question of whether our temporal map would be decipherable by deep-future humans or extra-terrestrial spacefaring civilizations is a deeply controversial one that intersects fields from philosophy of science to anthropology, semiotics, and cognitive science. As such, the question of its future intelligibility is outside the scope of our discussion here. To illustrate how future scientists and archaeologists may use the data provided in our time-map, however, we will assume that it will be intelligible by any spacefaring civilization; i.e., any civilization with the means to examine it.  \\subsection{Determining the Time Elapsed between Origin and Discovery}  Our Etching's temporal map contains a mixture of ``clocks,'' some of which are expected to run at constant rates over cosmological times (e.g., the hydrogen hyperfine transition frequency), while others' rates will vary (e.g., the Earth's spin period).  Depending  on the level of their scientific knowledge, the discoverers may be able to determine the Epoch of Origin of the Artifact by comparing Epoch of Origin  clock rates with Epoch of Discovery  clock rates.  For example, they could measure at the Epoch of Discovery, the periods $P$ and spindown rates   $\\dot{P}$ of some of the 13 MSPs whose Epoch of Origin periods are written on the etching. Then they could use Eq. 1 to determine $(T-T_0),$ the time elapsed between the two epochs.  If the elapsed time is sufficiently long, higher derivatives must be added to Eq. 1.  While current pulsar theory gives an expected value of the next derivative $\\ddot{P}  $ under the assumption of rotational energy loss to magnetic dipole radiation (Lyne \\& Graham-Smith 2012, p. 72), \\begin{equation} \\ddot{P}= - \\frac{\\dot{P}^2}  {P}, \\end{equation} this expression has been shown to be a poor approximation to reality in the few cases where the next derivative has actually been measured (Gradari et al 2011). Similar arguments apply to the other \ufffdclocks\ufffd such as the Earth's spin period and the Moon's orbital period.  Nevertheless, the discoverers should be  able to ascertain the true elapsed time by checking the consistency of their multiple results from an  ensemble of such ``clocks.''  \\subsection{Testing and Refining Scientific Theories}  Once the discoverers  establish $(T-T_0),$ the time elapsed since the origin of the Artifact, they  will also have the opportunity  to test and refine  some of their scientific theories over deep time.   For  example, they could combine our inscribed data with their measurements, thereby directly determining higher order derivatives as discussed above.  The assumptions underlying the theories could then be tested via comparison of these measured higher derivatives with the theoretically predicted ones, as shown in the above pulsar spindown example.  Some of our current clocks may ``break'' over cosmic timescales, in which case the discoverers will derive totally unique information by decoding our measurements.  For example, the pacing of resonances causing major changes in our spin axis obliquity is dictated by changes in tidal dissipation  in the Earth-Moon system, which is dependent on tectonic plate motions and other geophysical processes whose details are poorly known over long timescales, but which could be constrained by data from our era.  The discoverers  might further be able  to apply our currently recorded  measurements for uses that we cannot foresee, much as we now use ancient solar eclipse records to provide unique data on the evolution of the Earth's spin period over historic times (Stephenson 1997).  In addition, the mere discovery of the Artifact and spacecraft in geostationary space could provide the discoverers with a wealth of information on us, much as the discovery of the ``Antikythera Mechanism'' astronomical computer (Freeth et al 2006) in a shipwreck has done for our knowledge of ancient Mediterranean scientific and technological expertise. Further, we can hope that any collisions suffered by the Artifact before its discovery will not prevent its decoding, much as erosion of the Antikythera Mechanism slowed but did not stop current researchers from deciphering much of it."
    },
    "1208/1208.0825_arXiv.txt": {
        "abstract": "{Detailed studies of resolved young massive star clusters are necessary to determine their dynamical state and evaluate the importance of gas expulsion and early cluster evolution. In an effort to gain insight into the dynamical state of the young massive cluster R136 and obtain the first measurement of its velocity dispersion, we analyse multi-epoch spectroscopic data of the inner regions of 30 Doradus in the Large Magellanic Cloud (LMC) obtained as part of the VLT-FLAMES Tarantula Survey. Following a quantitative assessment of the variability, we use the radial velocities of non-variable sources to place an upper limit of 6\\,$\\kms$ on the line-of-sight velocity dispersion of stars within a projected distance of 5\\,pc from the centre of the cluster. After accounting for the contributions of undetected binaries and measurement errors through Monte Carlo simulations, we conclude that the true velocity dispersion is likely between 4 and 5\\,$\\kms$ given a range of standard assumptions about the binary distribution. This result is consistent with what is expected if the cluster is in virial equilibrium, suggesting that gas expulsion has not altered its dynamics. We find that the velocity dispersion would be $\\sim$25\\,$\\kms$ if binaries were not identified and rejected, confirming the importance of the multi-epoch strategy and the risk of interpreting velocity dispersion measurements of unresolved extragalactic young massive clusters.} ",
        "introduction": "The expression ``infant mortality\" of star clusters was initially coined by \\citet{lada2003} to describe the discrepancy between the number of observed open clusters and the number of embedded clusters. These authors argued that there are about ten times fewer open clusters than expected if all embedded clusters evolve into open clusters. The rapid removal of gas left-over from star formation was suggested to explain the apparent disruption of such a large fraction of clusters \\citep[e.g.][]{geyer, kroupaboily, BG2006}. The importance of the infant mortality scenario however depends on the definition adopted for embedded clusters \\citep{bressert2010, bastian2011}, with more conservative criteria requiring less than 50\\% of clusters to be destroyed to match the observed number of open clusters. But no matter which definition is adopted, the question of whether or not gas expulsion plays a significant role in the early evolution/disruption of star clusters still needs to be addressed.  Star clusters have been observed to expand in their first 100\\,Myr \\citep{2003MNRAS.338...85M, bastian2008, 2010ARA&A..48..431P}, but this expansion is not direct evidence for the importance of gas expulsion. There are two ways for clusters to expand as a response to mass loss \\citep[e.g.][]{hills1980}: (i) expansion following impulsive mass loss, e.g. change of potential due to removal of mass faster than the crossing time of the cluster, leaving the cluster in a super-virial state for a few crossing times, or (ii) adiabatic expansion, e.g. driven by stellar evolution on a slow timescale compared to the crossing time of the cluster \\citep[$\\sim$10\\,Myr for young massive clusters, e.g. ][]{2010ARA&A..48..431P}, in which case the cluster remains in virial equilibrium. Thus, the best way to evaluate the importance of rapid gas expulsion (case i) and the implications for the formation and early evolution of star clusters is to determine the dynamical state of young clusters. In particular, it is important to verify if clusters are in virial equilibrium from a young age.  Attempts to determine the dynamical state of young clusters have been made by comparing dynamical masses (obtained through measuring the velocity dispersion and size of a cluster) and photometric masses (estimated from the age and integrated luminosity). For several unresolved extragalactic star clusters with ages of less than $\\sim$10\\,Myr, the dynamical mass has been found to be up to ten times larger than the photometric mass \\citep[e.g.][]{bastianetal:2006}. This led to the suggestion that these clusters might be super-virial and expanding following gas expulsion \\citep{goodwinbastian:2006}. However, \\citet{gieles2010} showed that the increase in the measured line-of-sight velocity dispersion in these young clusters could be produced by the orbital motions of massive binaries. Massive stars indeed dominate the light of young massive clusters and their binary fraction is high \\citep[e.g.][Sana et al., submitted - hereafter Paper VIII]{mason2009, barba2010, sanaevans}. Therefore, the role of gas expulsion cannot be investigated through observations of unresolved extragalactic clusters.  Detailed dynamical studies of individual resolved clusters (i.e. in the Galaxy, Small Magellanic Cloud - SMC - or Large Magellanic Cloud - LMC) are necessary to make progress. In the case of radial velocity (RV) surveys, multi-epoch observations are needed to detect binaries, which can then be removed and allow a cleaner estimate of the velocity dispersion (unless the orbital solution is known, in which case the centre-of-mass velocity can be included in the dispersion calculation). The young massive cluster R136 \\citep[$M\\sim10^{5}$~$\\msun$,][]{2009ApJ...707.1347A}, in the 30~Doradus region of the LMC, is an ideal target to test the impact of gas expulsion given its young age of less than 2\\,Myr \\citep{koterheap, massey1998, Crowther2010}.  The identification of binaries was central to the study of stellar dynamics in 30 Doradus with the Gemini Multi-Object Spectrograph by \\citet{Bosch:2009}, who found that the velocity dispersion of NGC\\,2070 went from $\\sim$30\\,$\\kms$ to 8.3\\,$\\kms$ after correcting for orbital motions, confirming their prediction \\citep{bosch2001} that the high velocity dispersion of NGC 2070 could be due to undetected binaries. In order to perform a similar study for the dense surroundings of R136, a different observational approach was required because this region is too crowded for fibre spectroscopy. A key component of the VLT-FLAMES Tarantula Survey \\citep[VFTS;][hereafter Paper I]{evans:2011} is multi-epoch spectroscopy in the inner part of 30 Dor with the FLAMES--ARGUS integral-field unit (IFU), which we use here to obtain an estimate of the velocity dispersion of `single' stars in R136. The central velocity dispersion is an important proxy for the dynamical state of the cluster and is required to estimate, for example, the central potential and relaxation time-scale, both of which are required for detailed numerical ($N$-body) calculations of R136-like objects \\citep[e.g.][]{PZ:1999, gieles_mr:2010}. As a result of the same study, the identification of massive binaries could provide important clues to the star-formation process and the subsequent dynamical evolution of R136, in which binaries can be formed, ejected and destroyed. Our measurements will also serve as useful empirical input for the modeling of stellar interactions in dense clusters, such as the recent studies of \\citet{fujii} and \\citet{banerjee}.  We report here on a velocity dispersion estimate for the young massive cluster R136. We describe the FLAMES--ARGUS IFU data in Sect.~\\ref{argus_data}. In Sect.~\\ref{variability}, we present our RV and variability analysis of the stars observed with ARGUS. In Sect.~\\ref{class}, we discuss the spectral classification of the non-variable ARGUS sources.  In Sect.~\\ref{medusa}, we briefly introduce the VFTS Medusa data complementing the ARGUS data. We calculate the velocity dispersion from the stars showing no variability and also estimate the contribution of errors and undetected binaries in Sect.~\\ref{vdisp}. The implications of the measured velocity dispersion for the dynamical state of R136 are discussed in Sect.~\\ref{discuss}. Finally, we present our conclusions in Sect.~\\ref{conc}. ",
        "conclusions": "}  In an effort to determine the dynamical state of the young massive cluster R136, we used multi-epoch spectroscopy of stars in the inner regions of 30 Dor. We measured RVs with a Gaussian fitting procedure on selected key helium lines and performed a quantitative assessment of the variability. Out of 41 sources for which spectra were extracted from the ARGUS IFU data cubes, 16 were identified as non-variable. All of these were classified as O-type stars, three of which were also revealed to have composite spectra. To this sample of 16 ARGUS sources, we added measurements from 22 apparently single stars observed in the surrounding regions with Medusa-Giraffe, also as part of VFTS.  Using this sample of 38 non-variable massive stars within 10\\,pc from the centre of R136, we computed the velocity dispersion of the cluster. For the stars within 5\\,pc, we place an upper limit of 6\\,$\\kms$ on the line-of-sight velocity dispersion. This result does not change significantly if we exclude the few 3\\,$\\sigma$ outliers, supergiant candidates, and stars having composite spectra. We also noted that the measured velocity dispersion of the cluster includes a small contribution of $\\sim$0.5\\,$\\kms$ from rotation.  From Monte Carlo simulations, we established that the contribution of measurement errors to the observed velocity dispersion is almost negligible. We also estimated the contribution of undetected binaries, which is relatively small and dominated by long period systems beyond the detectability range of VFTS. When taking errors and undetected binaries into account, we estimate that the true velocity dispersion of the cluster (i.e. attributable only to cluster dynamics) is between 4 and 5\\,$\\kms$ for the stars within 5\\,pc from the centre.  Under basic assumptions, the expected central velocity dispersion in virial equilibrium was found to be  $\\approx$5.3\\,$\\kms$, in good agreement with our measurement, and we conclude that R136 is in virial equilibrium. Combined with the low velocity dispersions found in a few other young massive clusters, our results suggest that gas expulsion does not significantly alter the dynamics of these clusters.  We would have obtained a velocity dispersion of $\\sim$25\\,$\\kms$ if binaries had not been identified and rejected, which supports the suggestion that the alleged super-virial state of young star clusters can be explained by the orbital motions of binary stars \\citep{gieles2010}."
    },
    "1208/1208.2966_arXiv.txt": {
        "abstract": "We discuss the origin of two triggers of \\swift's Burst Alert Telescope (BAT) that occurred in 2011. The triggers were identified with \\new, a previously unknown X-ray source, and the known but unclassified X-ray transient \\source. We investigate the BAT data and follow-up observations obtained with \\swift's X-ray and ultraviolet/optical telescopes to demonstrate that both triggers are consistent with thermonuclear X-ray bursts. This implies that both sources are neutron star low-mass X-ray binaries. The total duration of $\\simeq7$~minutes and estimated energy output of $\\simeq(3-7)\\times10^{39}$~erg, fall in between that of normal and intermediately long X-ray bursts. From the observed peaks of the X-ray bursts, we estimate a distance of $\\lesssim 3.7$~kpc for \\new\\ and $\\lesssim 4.8$~kpc for \\source. We characterize the outburst and quiescent X-ray properties of the two sources. They have comparable average outburst luminosities of $\\simeq 10^{35-36}~\\lum$, and a quiescent luminosity equal to or lower than $\\simeq2\\times 10^{32}~\\lum$ (0.5--10 keV). \\new\\ returned to quiescence $\\simeq20$~d after its BAT trigger, while \\source\\ appears to exhibit long accretion outbursts that last several months to years. We identify a unique counterpart for \\source\\ in the ultraviolet/optical data. Finally, we serendipitously detect a flare lasting $\\simeq500$~s from an uncataloged X-ray/optical object that we tentatively classify as a flaring M-dwarf.\\\\ ",
        "introduction": "\\label{sec:intro} Low-mass X-ray binaries (LMXBs) consist of a neutron star or a black hole that accretes matter from a (sub-) solar mass companion star that overflows its Roche lobe. The accretion process involves the formation of an accretion disk and typically generates a 0.5--10 keV X-ray luminosity of $L_{\\mathrm{X}}\\simeq10^{36-39}~\\lum$ \\citep[e.g.,][]{chen97}, which places LMXBs amongst the brightest X-ray point sources in our Galaxy. Many LMXBs have an unstable accretion disk \\citep[see][for a review]{lasota01}. This causes the system to become \\textit{transient}, i.e., to alternate accretion outbursts with long episodes of quiescence during which the X-ray emission is much dimmer \\citep[$ \\lesssim10^{33}~\\lum$; e.g.,][]{menou99,garcia01}. Typically, the active phases of transient LMXBs last for a few weeks or months, while they reside in quiescence for years before a new outburst commences.  Whereas it is not straightforward to distinguish the nature of the compact object from the X-ray spectral and timing properties of the outburst \\citep[e.g.,][]{vanderklis1994,linares2007}, there are two phenomena that require a solid surface and are therefore considered a distinctive property of accreting neutron stars. These are coherent X-ray pulsations and thermonuclear X-ray bursts (i.e., type-I X-ray bursts, X-ray bursts hereafter). X-ray bursts are bright flashes of X-ray emission that may reach up to the Eddington limit and temporarily outshine the accretion luminosity. They are caused by unstable burning of accreted helium (He) and/or hydrogen (H) on the surface of the neutron star.  X-ray bursts are characterized by blackbody emission that peaks at a temperature up to $kT_{\\mathrm{bb}}\\simeq2-3$~keV and decays to $\\simeq 1$~keV as the X-ray burst fades. The observable properties such as the duration ($t_b$), radiated energy ($E_b$), and recurrence time ($t_{\\mathrm{rec}}$) depend on the amount and composition of the fuel that is consumed during the X-ray burst. The most commonly observed X-ray bursts last for $t_b \\simeq10-100$~s, have a radiated energy output of $E_b \\simeq 10^{39}$~erg, and recur on a timescale of minutes to days. To date, thousands of such events have been observed from about $\\simeq100$ different LMXBs \\citep[see, e.g.,][]{cornelisse2003,galloway06}.  Apart from these {\\it normal} X-ray bursts, there are two classes of more energetic events that are also more rare: intermediately long X-ray bursts \\citep[$t_b\\simeq$ tens of minutes to hours, $E_b\\simeq10^{40-41}$~erg, and $t_{\\mathrm{rec}}\\simeq$ weeks to years; see][]{chenevez2008,falanga08} and superbursts \\citep[$t_b\\simeq$ hours to days, $E_b\\simeq10^{42}$~erg, and $t_{\\mathrm{rec}}\\simeq$ a year; see][]{kuulkers2004,keek2008}. To date, only a few tens of such energetic X-ray bursts have been observed from about two dozens of LMXBs.  In this work, we discuss the properties of \\new\\ and \\source. We demonstrate that both sources displayed an X-ray burst in 2011 that triggered \\swift's Burst Alert Telescope \\citep[BAT;][]{barthelmy05}. We use rapid follow-up observations obtained with the narrow-field X-Ray Telescope \\citep[XRT;][]{burrows05} and Ultra-Violet/Optical Telescope \\citep[UVOT;][]{roming05} to fully characterize the trigger events. In addition, we investigate the outburst and the quiescent properties of both sources.     \\subsection{The New X-Ray Source \\new}\\label{subsec:new_intro} On 2011 June 24, \\swift/BAT triggered and located an unknown source \\citep[trigger 456014;][]{beardmore2011_gcn}. Based on the soft nature of the BAT detection and the proximity to the Galactic plane, the trigger was tentatively identified as a previously unknown Galactic transient and dubbed \\new\\ (J1850 hereafter). The detection of a fading X-ray counterpart in follow-up XRT observations provided an arcsecond localization of the source \\citep[][]{beardmore2011_gcn}. Based on the properties of the XRT data and the refined analysis of the BAT trigger data\\footnote{See also http://gcn.gsfc.nasa.gov/notices$\\_$s/456014/BA/}, it was suggested that this event was a thermonuclear X-ray burst \\citep[][]{markwardt2011,beardmore2011_atel}.  J1850 was not detected in the \\swift/BAT hard X-ray transient monitor on or after the time of the BAT trigger, with a 1-$\\sigma$ upper limit on the daily averaged 15--50 keV count rate of $7\\times10^{-4}~\\cnts~\\nh$ \\citep[][]{krimm2011}. However, the source was detected at an average 15--50 keV intensity of $(3.5\\pm1.0)\\times10^{-3}~\\cnts~\\nh$ (corresponding to a flux of $\\simeq2\\times10^{-10}~\\flux$ for a Crab-like spectrum) between 2011 May 18 and May 26. Archival searches back to 2005 February did not reveal any similar detections of J1850 \\citep[][]{krimm2011}.   \\begin{table} \\caption{Log of (Soft) X-ray Observations.} \\begin{center} \\begin{threeparttable} \\begin{tabular}{c c c c c} \\toprule Satellite/Instrument & Obs ID & Date & $t_{\\mathrm{exp}}$ (ks) \\\\ \\midrule \\multicolumn{4}{c}{{\\bf \\new}}  \\\\ \\xmm/EPIC* & 73740101 & 2003 Mar 21 & 26.8  \\\\ \\xmm/EPIC* & 73740201 & 2003 Mar 23 & 20.5   \\\\ \\xmm/EPIC* & 73740301 & 2003 Apr 18 & 24.2  \\\\ \\swift/XRT (PC+WT) & 456014000 & 2011 Jun 24 & 1.8\t \\\\ \\swift/XRT (PC+WT) & 456014001 & 2011 Jun 25 & 1.0\t \\\\ \\swift/XRT (PC+WT) & 456014002 & 2011 Jun 25 & 1.5\t \\\\ \\swift/XRT (PC+WT) & 456014003 & 2011 Jun 25 & 1.5\t \\\\ \\swift/XRT (WT) & 456014004 & 2011 Jun 26 & 1.0\t \\\\ \\swift/XRT (WT) & 456014005 & 2011 Jun 27 & 1.5\t \\\\ \\swift/XRT (WT) & 456014007 & 2011 Jun 28 & 4.9\t \\\\ \\swift/XRT (WT) & 456014008 & 2011 Jun 29 & 1.0\t \\\\ \\swift/XRT (WT) & 456014010 & 2011 Jul 1 & 1.0\t \\\\ \\swift/XRT (WT) & 456014011 & 2011 Jul 2 & 0.4\t \\\\ \\swift/XRT (WT) & 456014012 & 2011 Jul 3 & 1.2\t \\\\ \\swift/XRT (WT) & 456014013 & 2011 Jul 4 & 1.0 \t \\\\ \\swift/XRT (WT) & 456014014 & 2011 Jul 5 & 0.5\t \\\\ \\swift/XRT (WT) & 456014015 & 2011 Jul 6 & 1.0  \\\\ \\swift/XRT (WT) & 456014016 & 2011 Jul 7 & 1.0 \t \\\\ \\swift/XRT (WT) & 456014017 & 2011 Jul 8 & 0.9 \\\\ \\swift/XRT (WT) & 456014019 & 2011 Jul 10 & 0.5  \\\\ \\swift/XRT (PC) & 456014021 & 2011 Jul 12 & 1.1  \\\\ \\swift/XRT (PC) & 456014022 & 2011 Jul 13 & 1.0\t \\\\ \\swift/XRT (PC)* & 456014023 & 2011 Jul 16 & 0.2\t \\\\ \\swift/XRT (PC)* & 456014024 & 2011 Jul 17 & 0.8\t \\\\ \\swift/XRT (PC)* & 456014025 & 2011 Jul 18 & 0.8  \\\\ \\swift/XRT (PC)* & 456014026 & 2011 Jul 19 & 0.4  \\\\ \\swift/XRT (PC)* & 456014027 & 2011 Jul 20 & 0.1  \\\\ \\swift/XRT (PC)* & 456014028 & 2011 Jul 21 & 1.0  \\\\ \\midrule \\multicolumn{4}{c}{{\\bf \\source}}  \\\\ \\swift/XRT (PC+WT)$\\dagger$ & 35174001 & 2005 Jul 8 & 6.4 \\\\ % \\swift/XRT (PC+WT)$\\dagger$ & 35174003 & 2005 Oct 1 & 10.8  \\\\ % \\swift/XRT (PC+WT) & 35471001 & 2006 Mar 14 & 16.2  \\\\ % \\swift/XRT (PC+WT) & 35471002 & 2006 Jun 4 & 0.3  \\\\ % \\swift/XRT (PC+WT) & 35471003 & 2006 Jun 18 & 10.8  \\\\ % \\swift/XRT (PC)* & 35471004 & 2006 Oct 30 & 5.2 \\\\ % \\swift/XRT (PC)* & 35471005 & 2006 Nov 3 & 1.1 \\\\ % \\swift/XRT (PC)* & 35471006 & 2006 Nov 10 & 4.4 \\\\ % \\swift/XRT (PC)* & 35471007 & 2006 Nov 14 & 3.7 \\\\ % \\suzaku/XIS* & 702028010 & 2007 Apr 10 & 78.6 \\\\ % \\swift/XRT (PC+WT) & 35471008 & 2011 Aug 13 & 1.0 \\\\ % \\swift/XRT (PC+WT) & 35471009 & 2011 Aug 30 & 1.2 \\\\ % \\swift/XRT (PC+WT) & 35471010 & 2011 Aug 31 & 4.2 \\\\ % \\swift/XRT (WT) & 506913000 & 2011 Nov 3 & 1.9 \\\\ % \\swift/XRT (WT) & 35471011 & 2011 Nov 4 & 0.9 \\\\ % \\swift/XRT (WT) & 35471012 & 2011 Nov 5 & 1.0 \\\\ % \\swift/XRT (PC+WT) & 35471013 & 2012 Mar 9 & 0.7 \\\\ % \\swift/XRT (PC+WT) & 35471014 & 2012 Mar 14 & 1.0 \\\\ % \\swift/XRT (PC)* & 35471015 & 2012 May 25 & 1.2 \\\\ % \\swift/XRT (PC)* & 35471016 & 2012 Jun 9 & 1.9 \\\\ % \\swift/XRT (PC)* & 35471017 & 2012 Jun 11 & 2.0 \\\\ % \\swift/XRT (PC)* & 35471018 & 2012 Jun 13 & 0.9 \\\\ % \\swift/XRT (PC)* & 35471019 & 2012 Jun 15 & 0.7 \\\\ % \\swift/XRT (PC)* & 35471020 & 2012 Jun 17 & 0.3 \\\\ % \\swift/XRT (PC)* & 35471021 & 2012 Jun 21 & 2.3 \\\\ % \\bottomrule \\end{tabular} \\label{tab:obs} \\begin{tablenotes} \\item[]Note. -- Observations marked by an asterisk were obtained during quiescent episodes. The two observations marked by a dagger were also discussed by \\citet{falanga2006}. \\end{tablenotes} \\end{threeparttable} \\end{center} \\end{table}     \\subsection{The Unclassified Source \\source}\\label{subsec:source_intro} \\source\\ (hereafter J1922) is a transient X-ray source that was discovered during the \\swift/BAT hard X-ray survey of 2004 December -- 2005 March \\citep[][]{tueller2005a,tueller2005b}. \\swift/XRT follow-up observations allowed for the identification of a soft X-ray counterpart \\citep[][]{tueller2005b}. Non-detections with \\inte\\ \\citep[2003--2004;][]{falanga2006} and \\swift\\ \\citep[2006 October--November;][]{kennea2011_bursters} testified to its transient nature.  \\citet{falanga2006} discussed \\swift, \\rxte, and \\inte\\ data taken between 2005 July and November. The broadband spectrum and timing properties were found to be typical of an LMXB in a faint X-ray state, but the nature of the compact accretor (i.e., a neutron star or a black hole) could not be established. A combined blackbody and power-law model provided an adequate description of the broadband spectrum, with a hydrogen column density $N_{\\mathrm{H}} \\simeq (1.5-2.0)\\times10^{21}~\\nh$, spectral index $\\Gamma \\simeq 1.4-1.8$, and temperature $kT_{\\mathrm{bb}} \\simeq 0.4-0.6$~keV. The resulting 0.1--100 keV flux was $\\simeq (3.3-3.9)\\times10^{-10}~\\flux$ \\citep[][]{falanga2006}.  Renewed activity of J1922 was seen with \\maxi\\ and \\swift/XRT in 2011 August \\citep[][]{nakahira2011,kennea2011_bursters}. Inspection of the \\swift/BAT transient monitor data revealed that the source was detected in hard X-rays starting in 2011 July. It was seen at a mean 15--50 keV count rate of $(1.5 \\pm 0.3)\\times10^{-3}~\\cnts~\\mathrm{cm}^{-2}$ ($\\simeq1\\times10^{-10}~\\flux$ for a Crab-like spectrum) till early-August, after which the hard X-ray flux decreased \\citep[][]{kennea2011_bursters}. The 2011 intensity was similar to the average BAT count rate of the source in 2005.  \\swift/BAT triggered on J1922 on 2011 November 3 \\citep[][]{barthelmy2011}.\\footnote{See also http://gcn.gsfc.nasa.gov/notices$\\_$s/506913/BA/} Preliminary analysis of the BAT and XRT data revealed that the trigger was very likely caused by an X-ray burst \\citep[][]{degenaar2011_j1922}. Optical spectroscopy revealed clear He and H emission lines, supporting an LMXB nature \\citep[][]{wiersema2011,halpern2011}. The presence of H-lines implies that the donor star is H-rich and therefore the binary orbital period must be $P_{\\mathrm{orb}}\\gtrsim 1.5$ hr \\citep[][]{nelson86}. Upper limits on the quiescent optical counterpart ($>23.4$~mag in the $r$ and $g$ bands) are consistent with an M-dwarf or evolved star and suggest $P_{\\mathrm{orb}}\\lesssim5$~hr \\citep[][]{halpern2011}. ",
        "conclusions": "\\label{sec:discuss} We investigated three \\swift/BAT triggers that occurred in 2011. The BAT trigger spectra are soft and can be described by a blackbody model with a temperature of $kT_{\\mathrm{bb}}\\simeq 2-3$~keV. Rapid follow-up XRT observations were obtained for two of the triggers and revealed a decaying X-ray intensity that levels off to a constant value $\\simeq400$~s after the BAT peak. The X-ray tail shows a soft (thermal) emission spectrum that cools during the decay to $kT_{\\mathrm{bb}}\\simeq 0.6-0.7$~keV. Both events have a comparable exponential decay time ($\\tau \\simeq 85-100$~s), fluence ($f_{\\mathrm{total}}\\simeq2\\times10^{-6}~\\fluence$), and radiated energy output ($E_b\\simeq 3-7\\times10^{39}$~erg; Table~\\ref{tab:bursts}). These features are consistent with thermonuclear bursts occurring on accreting neutron stars, and strongly suggest that the BAT triggers were caused by type-I X-ray bursts.  We identified the X-ray transients J1850 and J1922 as the origin of the BAT triggers. This implies that both sources are neutron star LMXBs. J1850 is a previously unknown X-ray source, whereas J1922 was already discovered in 2005 but remained unclassified \\citep[][]{falanga2006}. We characterize the outburst and quiescent properties of these two new X-ray bursters.  The outburst of J1850 seen with \\swift's BAT and XRT had a duration of $\\simeq8$~weeks, but the onset of the outburst is not well constrained and hence it may have lasted (considerably) longer. The average 0.5--10 keV intensity of $L_{\\mathrm{X}}\\simeq 3\\times10^{35}~(D/3.7~\\mathrm{kpc})^2~\\lum$ classifies the source as a very-faint X-ray transient \\citep[Table~\\ref{tab:spec}; cf.][]{wijnands06}. At the time of the X-ray burst, the source was accreting at $\\simeq0.5\\%$ of the Eddington rate. Investigation of the ratio of XRT counts in different energy bands as the outburst progressed, revealed that the spectrum softens when the intensity decreases. This behavior has been seen in several neutron star and black hole LMXBs transitioning from the hard state to quiescence, although the underlying mechanism is not understood well \\citep[see][for a recent example and an overview]{armas2011}.  J1922 exhibited an outburst in 2005--2006 with an observed duration of $\\simeq20$~months and an 0.5--10 keV intensity of $L_{\\mathrm{X}}\\simeq 7\\times10^{35}~(D/4.8~\\mathrm{kpc})^2~\\lum$. The 2011--2012 outburst likely commenced in 2011 mid-July and had ceased by 2012 May, implying a duration of $\\simeq8-10$~months. The 0.5--10 keV luminosity of this second outburst was $L_{\\mathrm{X}}\\simeq1\\times10^{36}~(D/4.8~\\mathrm{kpc})^2~\\lum$. These outburst intensities fall in the very-faint to faint regime \\citep[Table~\\ref{tab:spec}; cf.][]{wijnands06}. At the time of the BAT trigger of 2011 November, the source was accreting at $\\simeq1.4\\%$ of the Eddington limit. Using \\swift/UVOT data, we have identified a unique UV/optical counterpart to J1922. This provides a sub-arcsecond localization of the source \\citep[][]{barthelmy2011}.  \\subsection{X-Ray Burst Properties of J1850 and J1922}\\label{subsec:long} The observable properties of X-ray bursts (e.g., the duration, recurrence time and radiated energy) depend on the conditions of the ignition layer, such as the thickness, temperature, and H-abundance. These can drastically change as the mass-accretion rate onto the neutron star varies, such that there exist distinct accretion regimes that give rise to X-ray bursts with different characteristics \\citep[][]{fujimoto81,bildsten98}. The overall similarities between the properties of the X-ray bursts of J1850 and J1922 (Table~\\ref{tab:bursts}) suggest that these events were ignited under similar conditions.  In the accretion regime of J1850 and J1922 ($\\simeq0.5\\%-1\\%$ of Eddington), the temperature of the burning layer is expected to be low enough for H to burn unstably. This may in turn trigger He ignition in an H-rich environment, which typically results in $\\simeq10-100$~s long bursts \\citep[e.g.,][]{fujimoto81,bildsten98,galloway06}. The detection of H-emission lines in the optical spectrum of J1922 suggests that the neutron star is accreting H-rich material \\citep[][]{wiersema2011,halpern2011}. This implies that the fuel triggering the X-ray burst could have indeed contained H. Given the similarities in burst properties, J1850 may therefore be expected to host an H-rich companion star as well.  The observed burst duration of $t_b\\simeq400$~s ($\\tau \\simeq 85-100$~s) is considerably longer than that typically observed for H-rich X-ray bursts \\citep[][]{chenevez2008,linares2012}. The burst profiles may look very different depending on the energy band in which they are observed \\citep[][]{chelovekov2006,linares2012}, and most bursts known to date have been detected with instruments that cover higher energies than \\swift/XRT \\citep[typically $>2$~keV, e.g., \\rxte, \\inte, \\beppo, \\fermi;][]{cornelisse2003,chenevez2008,galloway06,linares2012}. However, there are several X-ray bursts observed with \\swift/XRT, \\chan, or \\xmm\\ (i.e., in the same energy band as the bursts observed from J1850 and J1922) from neutron star LMXBs accreting at $\\simeq1\\%$ of Eddington that do show the expected duration of $\\simeq10-100$~s \\citep[e.g.,][]{boirin2007,trap09,degenaar09_gc,degenaar2012_gc}.  The uncommon properties of the bursts from J1850 and J1922 are further illustrated by their energetics. The estimated radiated energies of $E_b\\lesssim(3-7)\\times10^{39}$~erg are higher than that typically found for normal X-ray bursts, yet lower than those classified as intermediately long \\citep[][]{chenevez2008,linares2012}. There are several possible explanations that may account for the unusual burst properties.  First, for some LMXBs the first X-ray burst that occurred during a new outburst was found to be significantly longer than later bursts \\citep[e.g., Aql X-1 and Cen X-4;][]{fushiki1992,kuulkers2008}. It is thought that at the start of the outburst the neutron star is relatively cold and therefore a thicker layer of fuel can build up before the ignition conditions are met. Since the duration of an X-ray burst depends on the cooling time of the ignition layer, a thicker layer would result in a longer (and more energetic) X-ray burst. By considering the expected burst recurrence time, we can assess whether this scenario might be applicable to J1850 and J1922.  With the burst energetics at hand, the ignition depth can be estimated as $y=E_b (1+z)/4\\pi R^2 Q_{\\mathrm{nuc,burst}}$, where $z$ is gravitational redshift, $R$ is the neutron star radius and $Q_{\\mathrm{nuc,burst}}=1.6+4X~\\mathrm{MeV~nucleon}^{-1}$, the nuclear energy release during an X-ray burst given an H-fraction $X$ at ignition \\citep[e.g.,][]{galloway06}. For a neutron star with $M=1.4~\\Msun$ and $R=10$~km (i.e., $z=0.31$), we estimate $y\\simeq 10^{-8}~\\mathrm{g~cm}^{-2}$ for the bursts of J1850 and J1922. The recurrence time that corresponds to a given ignition depth is $t_{\\mathrm{rec}} \\simeq y(1+z)/\\dot{m}$, where $\\dot{m}$ is the accretion rate onto the neutron star surface per unit area. For J1850 and J1922 we infer an average accretion rate during outburst of $\\dot{m}=\\dot{M}/4\\pi R^2 \\simeq300$ and $\\simeq2300~\\mathrm{g~cm}^{-2}~\\mathrm{s}^{-1}$, respectively (assuming that the emission is isotropic). This would imply an expected recurrence time on the order of a few days to 2 weeks.  The BAT transient monitor revealed activity from J1850 $\\simeq5$ weeks before the BAT trigger of 2011 June 24. J1922 had been accreting for $\\simeq3$~months prior to its BAT trigger on 2011 November 3, as evidenced by \\swift\\ and \\maxi\\ (Sections~\\ref{subsec:new_intro} and~\\ref{subsec:source_intro}). Comparing this to the expected burst recurrence time of $\\lesssim2$ weeks suggests that the BAT triggers were likely not the first X-ray bursts that occurred during the outbursts of J1850 and J1922. The BAT coverage starts at 15~keV, which implies that it is not optimally sensitive to detect events as soft as X-ray bursts. These are therefore easily missed.  A second scenario that is worth considering is that some X-ray bursts display prolonged tails that last up to an hour \\citep[e.g., EXO 0748--676 and the ``clocked burster'' GS 1826--24;][]{zand09}. These are explained as the cooling of layers below the ignition layer, that were heated by inward conduction of energy generated during the X-ray burst. These long cooling tails are characterized by fluxes and fluences that are two orders of magnitude lower than the prompt burst emission. In the case of J1850 and J1922, however, the fluence in the BAT peak and XRT tail are of similar magnitude. This indicates that the long duration of these X-ray bursts is likely caused by a different mechanism.  Alternatively, the long burst duration may be explained in terms of a relatively large H-content in the ignition layer. The presence of H lengthens the duration of nuclear energy generation via the rp-process \\citep[][]{schatz2001}. There are two additional effects that likely contribute in making the bursts observable for a longer time. At low accretion rates the temperature in the neutron star envelop is expected to be low compared to brighter bursters. This implies that a thicker fuel layer can accumulate before reaching the critical ignition conditions, resulting in a longer burst. Furthermore, J1850 and J1922 accreted at a relatively low level of $\\simeq0.5\\%-1\\%$ of Eddington. For such low persistent emission levels, the tails of the X-ray bursts are visible for a longer time (if observed with a sensitive instrument), which can add to a longer burst duration. The burst recurrence time at the accretion rates inferred for J1850 and J1922 is considerably lower than for higher rates. This can account for the fact that most observed X-ray bursts are shorter than those seen for these two sources.  Intermediately long X-ray bursts are typically observed from sources accreting at $\\simeq0.1\\%-1\\%$ of Eddington. These events are both longer ($\\gtrsim10$~minutes) and more energetic ($\\simeq10^{40-41}$~erg) than we have observed for J1850 and J1922 \\citep[e.g.,][]{zand08,falanga09,linares09,degenaar2011_burst}. Some of these sources are strong candidate ultra-compact X-ray binaries. These contain an H-depleted companion star so that the neutron star accretes (nearly) pure He \\citep[e.g.,][]{zand08}. In absence of H-burning, the temperature in the accreted envelop is relatively low, so that a thick layer of He can build up before it eventually ignites in a long and energetic X-ray burst. A similarly long X-ray burst has also been observed, however, from a source that accretes at $\\simeq0.1$ of Eddington and shows a strong H-emission line in its optical spectrum. This testifies to the presence of an H-rich companion \\citep[][]{degenaar2010_burst}. In this case it may be that unstable H-burning could not immediately trigger He ignition, allowing the development of a thick layer of fuel \\citep[][]{cooper07,peng2007}.  Summarizing, the burst duration and energetics of the X-ray bursts observed from J1850 and J1922 appear to fall in between that of normal and intermediately long X-ray bursts \\citep[cf.][]{chenevez2008,linares2012}. Similar bursts have also been observed from the persistent neutron star LMXB 4U 0614+09 \\citep[][]{kuulkers09,linares2012}. We considered several possible scenarios that may account for the relatively long burst duration, and find it most likely that it is due to the ignition of a relatively thick layer of (H-rich) fuel. Our findings suggests that there may not exist two distinct groups of X-rays bursts, but rather a continuous range of burst durations and energies. This supports the idea proposed by \\citet{linares2012} that the apparent bimodal distribution is likely due to an observational bias toward detecting the longest and most energetic X-ray bursts from slowly accreting neutron stars.  \\subsection{Quiescent Properties of J1850 and J1922}\\label{subsec:q} J1850 could not be detected in archival \\xmm\\ observations with an estimated upper limit of $L_q \\lesssim(0.5-3.0)\\times10^{32}~(D/3.7~\\mathrm{kpc})^2~\\lum$. J1922 is detected in quiescence with \\swift/XRT and \\suzaku\\ at a 0.5--10 keV luminosity of $L_q \\simeq (0.4-1.0) \\times10^{32}~(D/4.8~\\mathrm{kpc})^2~\\lum$. These quiescent levels are common for neutron star LMXBs \\citep[e.g.,][]{menou99,garcia01,jonker2004}.  It is thought that the accretion of matter onto the surface of a neutron star compresses the stellar crust and induces a chain of nuclear reactions that deposit heat \\citep[][]{haensel1990a}. This heat spreads over the entire stellar body via thermal conduction and maintains the neutron star at a temperature that is set by the long-term averaged accretion rate of the binary \\citep[][]{brown1998,colpi2001}.  During quiescent episodes, the neutron star is  expected to thermally emit X-rays providing a candescent luminosity of $L_{\\mathrm{q,bol}} = \\langle \\dot{M} \\rangle Q_{\\mathrm{nuc}} / m_u$, where $Q_{\\mathrm{nuc}} \\simeq 2$ MeV is nuclear energy deposited in the crust per accreted baryon (\\citealt{haensel2008}; \\citealt{gupta07}, but see \\citealt{degenaar2011_terzan5_3}), $m_u=1.66\\times10^{-24}$~g is the atomic mass unit (i.e., $Q_{\\mathrm{nuc}} / m_u \\simeq 9.6 \\times 10^{17}~\\mathrm{erg~g}^{-1}$), and $\\langle \\dot{M} \\rangle$ is the long-term accretion rate of the binary averaged over $\\simeq10^4$~yr (i.e., including both outburst and quiescent episodes). The latter can be estimated by multiplying the average accretion rate observed during outburst ($\\dot{M}_{ob}$) with the duty cycle of the binary (i.e., the ratio of the outburst duration and recurrence time).  The outburst of J1850 observed with \\swift\\ in 2011 had a duration of $\\simeq8$~weeks (0.17~yr) and an estimated bolometric accretion luminosity of $L_{\\mathrm{bol}}\\simeq7\\times10^{35}~(D/3.7~\\mathrm{kpc})^2~\\lum$ (corresponding to $\\dot{M}_{\\mathrm{ob}}\\simeq6\\times10^{-11}~\\mdot$). The outburst duration and recurrence time are not known for J1850, but we can make an order of magnitude estimate based on the constraints of the quiescent luminosity.  Within the deep crustal heating model, a quiescent luminosity of $\\lesssim3\\times10^{32}~(D/3.7~\\mathrm{kpc})^2~\\lum$ would suggest a time-averaged mass-accretion rate of $\\langle \\dot{M} \\rangle \\lesssim 5 \\times 10^{-12}~\\mdot$. For $\\dot{M}_{\\mathrm{ob}}\\simeq6\\times10^{-11}~\\mdot$, the corresponding duty cycle is $\\lesssim10\\%$. If J1850  typically exhibits outbursts with a duration of 8 weeks, the expected recurrence time would be $\\gtrsim 2$~yr. We regard this as a lower limit, since the expected recurrence time increases for a longer outburst duration, or a quiescent thermal luminosity that is lower than the assumed upper limit of $3\\times10^{32}~\\lum$. This estimate is consistent with the fact that the \\swift/BAT hard transient monitor did not detect any other outbursts from J1850 back to 2005 February \\citep[][]{krimm2011}.  Since two outbursts with a relatively well-constrained duration have been observed for J1922, we can reverse the above reasoning and estimate the quiescent luminosity that is expected based on the observed duty cycle. The source was discovered when it exhibited an outburst of $\\simeq20$~months in 2005--2006. Renewed activity was observed from the source 5 yr later in 2011--2012, when it accreted for $\\simeq8-10$~months. For both outbursts, we estimate a similar bolometric accretion luminosity of $L_{\\mathrm{bol}}\\simeq(2-3)\\times10^{36}~\\lum$, which suggests a mean mass-accretion rate of $\\dot{M}_{\\mathrm{ob}}\\simeq3\\times10^{-10}~\\mdot$ (Table~\\ref{tab:spec}). If we assume a typical outburst duration of $\\simeq1.3$~yr and a recurrence time of $\\simeq5$~yr (i.e., a duty cycle of $\\simeq$26\\%), we can estimate a long-term averaged mass-accretion rate of $\\langle \\dot{M} \\rangle \\simeq 1\\times10^{-10}~\\mdot$ (equivalent to $\\simeq6\\times10^{15}~\\mathrm{g~s}^{-1}$).  If the outburst behavior observed over the past decade is typical for the long-term accretion history of J1922, then the estimated $\\langle \\dot{M} \\rangle$ should give rise to $L_{\\mathrm{q,bol}} \\simeq 6 \\times10^{33}~(D/4.8~\\mathrm{kpc})^2~\\lum$. This is considerably higher than the observed $L_q \\simeq 1 \\times10^{32}~(D/4.8~\\mathrm{kpc})^2~\\lum$ (0.5--10 keV). Although the bolometric luminosity may be a factor of a few higher than that measured in the 0.5--10 keV band, the luminosity remains lower than expected based on the crustal heating model. Limited by a low number of counts, we cannot constrain the shape of the quiescent spectrum of J1922 with the current available data. It is possible that the quiescent emission contains a significant non-thermal component \\citep[e.g.,][]{campana2005_amxps,wijnands05_amxps,heinke2009,cackett2010_cenx4, degenaar2012_1745,degenaar2012_amxp}, which would increase the discrepancy between the expected and observed quiescent thermal emission.  Provided that the heating models are correct, a plausible explanation for this apparent mismatch could be that the time-averaged mass-accretion rate is lower than estimated (e.g., because the recent outburst behavior is not representative for the long-term accretion history). Alternatively, the neutron star could be cooling faster than assumed in the standard paradigm \\citep[][]{page2004}. A comparison with theoretical cooling models of \\citet{yakovlev2004} would suggest enhanced cooling due to the presence of kaons in the neutron star core.  Four neutron star LMXBs have been closely monitored after the cessation of their very long ($\\gtrsim1$~yr) accretion outbursts, which has revealed that the neutron star temperature was gradually decreasing over the course of several years \\citep[][]{wijnands2001,wijnands2003,cackett2008,cackett2010,degenaar09_exo1,degenaar2010_exo2,fridriksson2010,fridriksson2011,diaztrigo2011}. Recently, similar behavior has been observed for a transient neutron star LMXB in Terzan 5 that exhibited a much shorter accretion outburst of $\\simeq10$~weeks \\citep[][]{degenaar2011_terzan5_2,degenaar2011_terzan5_3}. These observations can be explained as cooling of the neutron star crust, which became considerably heated during the accretion outburst and needs time to cool down in quiescence. Monitoring and modeling this crustal cooling provides the unique opportunity to gain insight into the properties of the neutron star crust and core \\citep[][]{rutledge2002,wijnands04_quasip,shternin07,brown08,degenaar2011_terzan5_3,page2012}.  With its relatively low quiescent luminosity, long outburst duration and low extinction, J1922 would be a promising target to search for crustal cooling now that its recent outburst has ceased."
    },
    "1208/1208.0152_arXiv.txt": {
        "abstract": "{Surface abundances of OB stars provide important information on  mixing processes in stellar interiors and on evolutionary status. Performing abundance determination from selected C, N and O lines in the optical spectra of OB stars is thus necessary to understand the evolution of massive stars. } {In this study, we investigate the formation of \\ion{C}{iii} 4647--51--50 and \\ciii\\ in the atmosphere of O stars to see if they can be reliably used for abundance determinations.} {We use atmosphere models computed with the code CMFGEN. The key physical ingredients explaining the formation of the \\ion{C}{iii} lines are extracted from comparisons of models with different stellar parameters and through examining rates controlling the level populations.} {The strength of \\ciii\\ critically depends on UV \\ion{C}{iii} lines at 386, 574 and 884 \\AA. These lines control the \\ciii\\ upper and lower level population. \\ion{C}{iii} 884 plays a key role in late O stars where it drains the lower level of \\ciii. \\ion{C}{iii} 386 and \\ion{C}{iii} 574 are more important at early spectral types. The overlap of these UV lines with \\ion{Fe}{iv} 386.262, \\ion{Fe}{iv} 574.232 and \\ion{S}{v} 884.531 influences the radiative transfer at 386, 574 and 884 \\AA, and consequently affects the strength of \\ciii. \\ion{C}{iii} 4650 is mainly controlled by the \\ion{C}{iii} 538 line which acts as a drain on its lower level. \\ion{Fe}{iv} 538.057 interacts with \\ion{C}{iii} 538 and has an impact on the \\ion{C}{iii} 4650 profile. Low temperature dielectronic recombinations have a negligible effect on the line profiles. The main effect of the wind on the \\ion{C}{iii} lines is an increase of the emission as mass loss rate increases, due to a shift of the formation region towards high velocity. Such a shift also implies desaturation of \\ion{C}{iii} 386,  \\ion{C}{iii} 538,  \\ion{C}{iii} 574 and \\ion{C}{iii} 884 which in turn modify the \\ion{C}{iii} 4650 and \\ciii\\ line profiles. Metallicity effects are complex and result from an interplay between ionization, metallic line interaction with UV \\ion{C}{iii} lines, and wind effects. } {Given our current understanding of the stellar and wind properties of O stars, and in view of the present results, the determination of accurate carbon abundances from  \\ion{C}{iii} 4647--51--50 and \\ciii\\ is an extremely challenging task. Uncertainties lower than a factor of two on C/H determinations based only on these two sets of lines should be regarder as highly doubtful. Our results provide a possible explanation of the variability of \\ion{C}{iii} 4650 in Of?p stars.} ",
        "introduction": "\\label{s_intro}  Massive stars turn hydrogen into helium through the CNO cycle in their core. In this process, nitrogen is produced at the expense of carbon and oxygen. Mixing processes subsequently transport the products of nucleosynthesis from the core to the outer layers and surface of the star. The spectrum of a massive star thus contains the imprints of nuclear reactions taking place in the interior of the star. The older a star, the larger its nitrogen content and the lower its carbon abundance (provided the star is of type O and not yet an evolved object such as Wolf-Rayet (W-R), red supergiant or Luminous Blue Variable). The surface chemical composition depends on the star's rotation rate, metallicity, and mass. A fast rotation, for example, favors the transport of chemical species and thus quicker enrichment of the photosphere by material synthesized in the core \\citep{mm00}. Low metallicity stars burn hydrogen more efficiently, as does a 60 \\msun\\ stars compared to a 20 \\msun\\ star. Chemical surface abundances are thus a key to understand the physical processes controlling the evolution of massive stars. It was surface abundances that provided key clues that W-R stars were evolved \\citep[e.g.,][]{S73_Habund,SW82_CNO}.  Recent results \\citep{hunter08,przy10,brott11a,martins12} have shown that strong constraints on the physics of massive stars could be obtained by spectroscopic analysis and abundance studies.  Nevertheless, the determination of surface abundances of O stars is challenging. Non--LTE atmosphere models including line--blanketing and winds are required, and there is a rarity of metallic lines in O star spectra. The UV range contains  a few strong resonance lines (\\ion{C}{iii} 1176, \\ion{N}{v} 1240, \\ion{C}{iv} 1550, \\ion{N}{iv} 1720) but they are mainly formed in the wind part of the atmosphere where numerous effects (X-rays, clumping) render their modelling difficult. Photospheric lines are found in the optical range, but, here again, only a few lines are available.  \\citet{rgonz11a} have revisited the formation mechanism of \\ion{N}{iii} 4634--4640 and have used several \\ion{N}{ii}, \\ion{N}{iii} and \\ion{N}{IV} lines to derive abundances of O stars in the Magellanic Clouds. \\citet{martins12} studied the effect of magnetism on the surface abundance of O stars, and based their nitrogen abundance measurements on \\ion{N}{ii}, \\ion{N}{iii} and \\ion{N}{iv} optical lines. Nitrogen is a crucial indicator of chemical enrichment. But more precise constraints can be brought to stellar evolution if the ratio N/C can be derived because N/C is less sensitive to the initial metal content than N/H. For instance the ratio of N/C in the Galaxy to N/C in the SMC is less that 3, while N/H in the Galaxy is 13 times larger than N/H in the SMC. Thus N/C, especially when combined to the N/O ratio \\citep[e.g.][]{przy10}, provides stronger constraints on the enrichment and mixing history of CNO material at a given metallicity.  Unfortunately, carbon lines are even less numerous than nitrogen lines in O stars spectra. In the optical range, the strongest lines are  \\ion{C}{iii} 4647--50--51 (hereafter \\ion{C}{iii} 4650), \\ciii, and \\ion{C}{iv} 5801--12. \\ion{C}{iii} 4650 is observed in absorption in most O stars. It shows emission only in the earliest O supergiants. In early O dwarfs, it is almost absent. A peculiar class of objects -- the Ofc stars \\citep{walborn10} -- displays always \\ion{C}{iii} 4650 emission. Stars of the Ofc class encompass Of?p stars which, among other characteristics, have a strong but varying \\ion{C}{iii} 4650 emission. Many of them are magnetic \\citep{donati06,hubrig08,martins10,wade11}. Ofc stars are found mainly at spectral type O5 but at all luminosity classes \\citep{walborn10}. So far, no explanation has been given for the appearance of such \\ion{C}{iii} 4650 emission. Its analysis is rendered difficult by the presence of the neighbouring strong \\ion{N}{iii} 4634--4640 lines. \\ion{C}{iv} 5801--12 is present in most O stars, but is often not observed. In late-type stars it is in absorption, while in early type stars it can exhibit a complex profile with broad emission wings and absorption \\citep{conti74,jc12}.  \\ciii\\ is present in most O stars except the earliest ones \\citep{walborn80}. It is in emission in giants and supergiants, and in most dwarfs. The emission is on average stronger in supergiants than in dwarfs \\citep{conti74}. Contrary to \\ion{C}{iii} 4650, \\ciii\\ is located in a spectral region free of other lines and hence its profile is well defined. Spectroscopic analysis have generally not considered this line for carbon abundance determination -- atmosphere models often have difficulties reproducing the observed profiles, even if the H, \\ion{He}{i}, \\ion{He}{ii} and UV CNO lines are well fitted. This indicates that the formation of that line is complex and probably depends on missing or improperly accounted for physics. \\citet{nu71} found that the \\ciii\\ emission was due to a depopulation of the lower level by \\ion{C}{iii} 884. Hence, a correct description of the radiative transfer close to 884 \\AA\\ is required. Since the models used by \\citet{nu71} did not include a proper line--blanketing and non--LTE treatment, we revisit the formation of \\ciii. We also investigate the formation of \\ion{C}{iii} 4650.  In Section \\ref{s_mod} we present the method and models we have used. Sect.\\ \\ref{form_sing} describes the formation mechanism of \\ciii\\ and Sect.\\ \\ref{form_trip} focuses on \\ion{C}{iii} 4650. The wind effects are gathered in Sect.\\ \\ref{s_wind}. We summarize our results in Sect.\\ \\ref{s_conc}. ",
        "conclusions": "\\label{s_conc}  We have investigated the formation processes of \\ciii\\ and \\ion{C}{iii} 4650 in the atmosphere of O stars. We have used the non--LTE atmosphere code CMFGEN to perform simulations at two effective temperatures (30500 K and 42000 K) corresponding to late and early spectral types. Our results can be summarized as follows:  \\begin{itemize}  \\item[$\\bullet$] For wind densities typical of O stars, the formation of \\ciii\\ is controlled by several UV transitions: \\ion{C}{iii} 386, \\ion{C}{iii} 574 and \\ion{C}{iii} 884. The latter is crucial at low \\teff: when it is removed from models, \\ciii\\ switches from emission to absorption. \\ion{C}{iii} 574 and \\ion{C}{iii} 386 play the same role at high \\teff.  \\item[$\\bullet$] \\ion{C}{iii} 884 acts as a drain on the \\lsl\\ population (lower level of \\ciii). When it is removed, \\lsl\\ is more populated and emission is reduced/suppressed. \\ion{C}{iii} 574 drains the \\lsu\\ population and \\ion{C}{iii} 386 populates \\lsl. When they are removed, \\lsu\\ is more populated and \\lsl\\ is depopulated, favoring emission.  \\item[$\\bullet$] Lower gravity models have larger ionizing flux. As a consequence, the lower level of \\ion{C}{iii} 884 (2p2p~$^1$F$^{\\scriptsize \\rm o}$), which is collisionally coupled to the ground level, is depopulated at low \\logg. The drain of \\lsl\\ by \\ion{C}{iii} 884 is thus increased. This leads to an increase of \\ciii\\ emission when \\logg\\ decreases at 30500 K. At \\teff\\ = 42000 K, the effect of \\logg\\ is opposite and is due to the role of \\ion{C}{iii} 386 in populating the lower level of \\ciii.  \\item[$\\bullet$] The strength of \\ciii\\ is very sensitive to the action of a few metallic lines on the radiative transfer near 386, 574 and 884 \\AA. In particular, \\ion{Fe}{iv} 386.262 and \\ion{Fe}{iv} 574.232, when removed from models, lead to a stronger \\ciii\\ emission. If \\ion{S}{v} 884.531 is removed, the \\ciii\\ emission is reduced. The reason is the overlap of these features with the key \\ion{C}{iii} lines controlling the shape of \\ciii. As an example, \\ion{Fe}{iv} 574.232 extracts photons from \\ion{C}{iii} 574 leading to a lower population of \\lsu. Consequently, \\ion{Fe}{iv} 574.232 limits the \\ciii\\ emission. When it is removed, \\lsu\\ is more populated and the \\ciii\\ emission is stronger.  \\item[$\\bullet$] The formation of \\ion{C}{iii} 4650 is controlled by radiative transfer at 538 \\AA. The \\ion{C}{iii} 538 line drains the population of \\ltl, the lower level of \\ion{C}{iii} 4650. When \\ion{C}{iii} 538 is removed, \\ltl\\ is more populated and \\ion{C}{iii} 4650 switches from emission to absorption, or shows a stronger absorption.  \\item[$\\bullet$] \\ion{Fe}{iv} 538.057 enhances photon escape in the \\ion{C}{iii} 538 line which helps to reduce the population of \\ltl, and thus weakens the \\ion{C}{iii} 4650 absorption (or enhances the emission).  \\item[$\\bullet$] Low temperature dielectronic recombinations have a limited effect on the appearance of \\ciii\\ and \\ion{C}{iii} 4650 in O stars.  \\item[$\\bullet$] When the mass loss rate increases, a growing fraction of the \\ion{C}{iii} 4650 and \\ciii\\ lines is formed in the wind, leading to a larger emission. In addition, as \\mdot\\ gets higher, the formation region of the \\ion{C}{iii} lines (UV and optical) shifts to higher velocities. The associated Doppler shifts lead to desaturation of the key \\ion{C}{iii} UV lines. This affects the level populations and consequently the intrinsic \\ion{C}{iii} 4650 and \\ciii\\ line profiles. At low mass-loss rates, the change in line profile with a change in the mass-loss rate is not monotonic.  \\item[$\\bullet$] Metallicity effects are a complex interplay between \\ion{C}{iii} UV lines, metallic \\ion{Fe}{iv} and \\ion{S}{v} lines, ionization, winds and abundance effects. There is no simple way of predicting how \\ion{C}{iii} 4650 and \\ciii\\ will react to a change of metallicity or abundance of specific elements.  \\end{itemize}  In view of our results, it seems highly risky to derive accurate carbon abundances from \\ion{C}{iii} 4650 and \\ciii\\ in O star atmospheres. Uncertainties of at least a factor of two on C/H are expected just because of the uncertainty on atomic data and wavelengths of metallic lines. Priority should be given to other carbon lines, such as the \\ion{C}{iii} 2s4f\\,$^3$F$^{\\scriptsize \\rm o}$-2s5g\\,$^3$G multiplet at 4070\\,\\AA, or even UV lines (\\ion{C}{iv} 1164 and \\ion{C}{iii} 1176) provided the wind is not too strong and they are mainly photospheric.  The current results can be used to try to understand the behavior of \\ion{C}{iii} 4650 in Of?p stars, many (all?) of which are magnetic  \\citep{donati06,hubrig08,martins10,wade11,wade191612}. Of?p stars show periodic variability, especially in \\ion{C}{iii} 4650 emission which can be as strong as in the neighbouring \\ion{N}{iii} lines before almost vanishing. H$_{\\alpha}$ shows the same variability and is strongest when \\ion{C}{iii} is at its maximum emission strength. \\citet{sund12} demonstrated that the changes in H$_{\\alpha}$ emission of HD~191612 could be understood as a result of viewing the dynamic magnetosphere from different angles.  Due to wind confinement by the magnetic field, an overdense region is created in the magnetic equator. If the star is seen from the magnetic pole, a large surface of the overdense region is observed and the H$_{\\alpha}$ emission is stronger than if the star is seen from the magnetic equator. In the present study, we have seen that \\ion{C}{iii} 4650 can show emission if the wind contribution becomes dominant, and/or if for some reason the drain effect of \\ion{C}{iii} 538 is increased so that \\ltl\\ is depopulated. If we observe HD~191612 from the magnetic pole, the wind velocity at a given radius is larger than in the magnetic equator, due to confinement (see Fig.\\ 2 of Sundqvist et al.\\ 2012). The role of the wind is thus stronger, favoring \\ion{C}{iii} 4650 emission. In addition, since the atmosphere density is lower in the magnetic pole direction, we can expect a larger ionizing flux due to reduced opacity. As seen in our study, this favors the drain of \\ltl\\ through \\ion{C}{iii} 538, and thus \\ion{C}{iii} 4650 emission. Hence, the existence of an equatorial overdensity could lead to the observed \\ion{C}{iii} 4650: more emission when the star is seen pole--on, less emission when it is seen equator--on, just as H$_{\\alpha}$. Of course, this needs to be tested by detailed 2D radiative transfer calculation, but the present results provide a good working frame."
    },
    "1208/1208.0478_arXiv.txt": {
        "abstract": "Probabilities of a pair of fermions and  bosons  creation  in a static and spatially  uniform electric field $E$ are represented in the Schwinger formulas by infinite series. It is believed that in weak fields the main contribution to the probability is given by the first term of series, however the  size of the remainder apparently was analyzed by nobody. We study the mathematical structure of the Schwinger series by using methods developed during last decades and  prove that the given series allows an  exact summation and the contribution of remainder growths rapidly with the field strength. As a rule, it is argued that the pair of particles begin to be produced efficiently from the vacuum only in the fields of strength  $E\\geq E_{cr}$. However, the direct calculation shows  that the Schwinger formula for creation of $e^+e^-$ pairs  is valid only  at the field intensities of $E<0.0291E_{cr}$.  At  higher fields, the probability of pair production in an unit space-time volume exceeds unity.  In this regard, we refine the formula for the probability of pair creation and numerically  find that in the field of strength $2.95\\%$ of $E_{cr}$ the  pair production probability  is almost 100$\\%$. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.5887_arXiv.txt": {
        "abstract": "Magnetic diffusion in accretion flows changes the structure and angular momentum of the accreting material. We present two power law similarity solutions for flattened accretion flows in the presence of magnetic diffusion: a secularly-evolving Keplerian disc and a magnetically-diluted free fall onto the central object. The influence of Hall diffusion on the solutions is evident even when this is small compared to ambipolar and Ohmic diffusion, as the surface density, accretion rate and angular momentum in the flow all depend upon the product $\\eta_H(\\mathbf{B}\\cdot\\Omega)$, and the inclusion of Hall diffusion may be the solution to the magnetic braking catastrophe of star formation simulations. ",
        "introduction": "Accretion flows onto protostars are flattened and braked by magnetic fields. The transportation of angular momentum from the flow to the envelope is facilitated by the bending of the magnetic field lines and inhibited by the drift of those lines against the flow\\footnote{This is only strictly true for ambipolar and Ohmic diffusion.}. In magnetohydrodynamical (MHD) simulations of star formation it has been shown that it is possible to remove all of the angular momentum from the accretion flow using a weak field, with the consequence that no accretion disc forms \\citep[e.g.][]{pb2007, ml2008}. Magnetic diffusion reduces the braking so that a protostellar disc forms \\citep[e.g.][hereafter KK02; Mellon \\& Li 2009]{kk2002}; these discs are precursors to those in which planet formation occurs.  The amount of magnetic braking affecting an accretion flow depends on the coupling of the largely neutral medium to the magnetic field, which is facilitated by collisions between the neutral and charged particles that transmit the Lorentz force to neutrals. In most simulations the magnetic field behaviour has been approximated by ideal MHD (where the magnetic field is frozen into the neutral particles), however this approximation breaks down as the density in the disc increases. The relative drifts of different charged species with respect to the neutral particles delineate three conductivity regimes. \\textit{Ohmic (resistive)} and \\textit{ambipolar diffusion}, which behave similarly in a thin disc (to a first order approximation), have been shown to reduce the effectiveness of magnetic braking \\citep{sglc2006, mim2007, ml2009}. \\textit{Hall diffusion} occurs when the degree of coupling between the different charged species and the neutrals varies, and is expected to dominate in large regions in accretion flows \\citep{ss2002, w2007}. For certain field configurations Hall diffusion can in principle cause the magnetic field to be accreted faster than the neutral fluid \\citep{wn1999, bw2011}.  The role of Hall diffusion in star formation and accretion discs is neither fully understood nor explored due to the difficulty of performing numerical simulations \\citep[although this is starting to change, e.g.][]{lks2011, kls2011}, highlighting the need for simpler analytic models that demonstrate the importance of the Hall effect. In this paper we present two power law similarity solutions to the MHD equations for an isothermal thin disc with Hall, ambipolar and Ohmic diffusion: one in which the disc is Keplerian and a second in which the collapsing material spirals onto the central object without disc formation. We discuss the implications and the physics controlling these remarkable solutions.  \\section[]{Formulation}\\label{equations}  The magnetohydrodynamic equations for an isothermal system are given by \\begin{equation} \\frac{\\partial\\rho}{\\partial{t}} + \\nabla\\cdot(\\rho\\mathbf{V}) = 0, \\label{m1} \\end{equation} \\begin{align} \\rho\\frac{\\partial\\mathbf{V}}{\\partial{t}} + \\rho(\\mathbf{V}\\cdot\\nabla)\\mathbf{V} = -\\nabla{P} - \\rho\\nabla\\Phi + \\mathbf{J}\\times\\mathbf{B}, \\label{rm1} \\end{align}  \\begin{align} \\nabla^2\\Phi = {4{\\pi}G\\rho}, \\label{pot1} \\end{align} \\begin{align} \\nabla\\cdot\\mathbf{B}=0, \\label{am1} \\end{align} and \\begin{align} \\frac{\\partial{\\mathbf{B}}}{\\partial{t}} &= \\nabla \\times (\\mathbf{V} \\times \\mathbf{B}) \\nonumber \\\\ &-\\!\\nabla\\!\\times\\! \\left[\\eta\\!\\left(\\nabla\\!\\times\\!\\mathbf{B}\\right) + \\eta_{H}\\!\\left(\\nabla\\!\\times\\!\\mathbf{B}\\right)\\!\\times\\!\\mathbf{\\hat{B}} + \\eta_{A}\\!\\left(\\nabla\\!\\times\\!\\mathbf{B}\\right)_{\\perp}\\!\\right]\\!, \\label{in1} \\end{align} where $\\rho$ is the gas density, $\\mathbf{V}$ the velocity field, $P$ the midplane gas pressure, $\\Phi$ the gravitational potential (defined as $\\mathbf{g} = -\\nabla\\Phi$ where $\\mathbf{g}$ is the gravitational field), $G$ the gravitational constant, $\\mathbf{B}$ the magnetic field and $\\mathbf{J}$ the current density defined as $\\mathbf{J} = c(\\nabla \\times \\mathbf{B})/4\\pi$ where $c$ is the speed of light. The diffusion coefficients for the Ohmic, Hall and ambipolar terms in the induction equation are $\\eta$, $\\eta_H$ and $\\eta_A$ respectively.  We adopt cylindrical coordinates and assume that the disc is axisymmetric and thin. The velocity field, $\\mathbf{V}(r)$, has both radial and azimuthal field components, the latter giving rise to the specific angular momentum $J = r V_\\phi$, as the material is assumed to settle rapidly to the disc midplane. We assume that the disc is threaded by an open magnetic field that is symmetric about the midplane, defining $B_z$ as the vertical field component at the midplane and $B_{rs}$ and $B_{\\phi s}$ as the components at the surface $z = H$, where $H(r)$ is the half-thickness of the disc. We neglect any mass loss due to a disc wind. As the disc is isothermal, the pressure at the midplane is given by $P = \\rho c_s^2$, where $c_s$ is the constant isothermal sound speed, typically taken to be $0.19$ km s$^{-1}$.  Equations \\ref{m1}--\\ref{in1} are vertically-averaged as in \\citet[following KK02]{bw2011} by assuming the disc is thin and integrating over $z$, which allows us to discard any terms of order $H/r$. The density, radial velocity, azimuthal velocity and radial gravity are approximated as being constant with height, as are the quantities $\\eta$, $\\eta_H/B$ and $\\eta_A/B^2$. The radial and azimuthal magnetic field components are taken to scale linearly with height, so that their values at the disc surface ($B_{rs}$ and $B_{\\phi s}$) may be used to characterise the field.  The equations are further simplified by employing monopole expressions for $B_{rs}$ and $g_r$: \\begin{align} B_{rs} &= \\frac{\\Psi(r,t)}{2\\pi{}r^2} \\label{b_rs1}\\\\ \\text{and }g_r &= -\\frac{GM(r,t)}{r^2},\u00e5 \\label{g_rs1} \\end{align} where $M$ and $\\Psi$ are the mass and magnetic flux enclosed within a radius $r$. These describe the disc sufficiently well that a more complicated iterative method of calculating the field and gravity is unnecessary \\citep{cck1998}.  The azimuthal field component is calculated by balancing the torques on the disc from rotation and magnetic braking, which transports angular momentum from the disc to the low-density external medium into which the field is frozen \\citep{bm1994}. Assuming that the background angular frequency is small and that the external Alfv\\'en wave speed is constant and parameterised by $V_\\text{A,ext} = c_s/\\alpha$, the azimuthal field component is then \\begin{align} B_{\\phi s}=-\\mathrm{min}\\Biggl[\\frac{\\Psi\\alpha}{\\pi{r^2}c_s} &\\left[\\frac{J}{r}-\\frac{\\eta_HB_{rs}}{HB}\\right]\\nonumber\\\\ &\\left[1+\\frac{\\Psi\\alpha}{\\pi{r^2}c_s}\\frac{\\eta_A}{B^2} \\frac{B_z}{H}\\right]^{-1};\\delta{B_z}\\Biggr] \\label{b_phisfinal} \\end{align} \\citep{bw2011}. The field is capped as a way of representing the many magnetohydrodynamical instabilities and processes that might prevent the azimuthal field component from exceeding the poloidal component such as internal kinks, turbulence and the magnetorotational instability (MRI). The cap $\\delta = 1$ corresponds to the typical value of $|B_{\\phi s}|/B_z$ obtained when the vertical angular momentum transport is dominated by a centrifugally-driven disc wind (KK02).  The vertically-averaged equations are then: \\begin{align} &\\frac{\\partial\\Sigma}{\\partial{t}} + \\frac{1}{r}\\frac{\\partial}{\\partial{r}}(r\\Sigma{V_{r}}) = 0, \\label{mass1}\\\\ \\nonumber\\\\ &\\frac{\\partial{V_{r}}}{\\partial{t}} + {V_{r}}\\frac{\\partial{V_{r}}}{\\partial{r}} = g_r - \\frac{c_s^{2}}{\\Sigma}\\frac{\\partial\\Sigma}{\\partial{r}} + \\frac{B_{z}B_{rs}}{2\\pi\\Sigma} + \\frac{J^2}{r^3}, \\label{rad1}\\\\ \\nonumber\\\\ &\\frac{\\partial{J}}{\\partial{t}} + V_{r}\\frac{\\partial{J}}{\\partial{r}} = \\frac{rB_{z}B_{\\phi s}}{2\\pi\\Sigma},\\label{ang1}\\\\ \\nonumber\\\\ &\\frac{{\\Sigma}c_s^2}{2H} = \\frac{\\pi}{2}G\\Sigma^2 + \\frac{GM_{c}\\Sigma{H}}{4r^3} + \\frac{1}{8\\pi}\\left(B_{rs}^{2} + B_{\\phi s}^2\\right),\\label{vert1} \\end{align} and \\begin{align} &\\frac{H}{2\\pi}\\frac{\\partial\\Psi}{\\partial{t}} = -rHV_rB_z - \\eta{B_{rs}} - \\frac{r\\eta_H}{B}B_zB_{\\phi s} - \\frac{r\\eta_A}{B^2}B_{rs}B_z^2; \\label{ind1} \\end{align} $\\Sigma$ is the column density, defined as $\\Sigma = 2H\\rho$, and $M_c$ is the mass of the central star. These equations are a simplified set of those used in \\citet{bw2011}, removing the $H\\partial{B_z}/\\partial{r}$ terms that were used to refine the structure of the disc and shocks in the full similarity solutions.  The Ohmic and ambipolar diffusion terms scale together, to a zeroth-order approximation, as they possess a similar dependence upon $B$ and appear in the induction equation multiplied with the same field component. Ohmic diffusion is expected to be more important in the inner regions where the density is high, while ambipolar diffusion dominates in the outer regions \\citep{w2007}. As the field within the disc is effectively vertical, ambipolar and Ohmic diffusion influence the field drift in the same direction, and only one term is required to study the change in disc behaviour introduced by Hall diffusion. We then describe ambipolar and Ohmic diffusion using the Pedersen diffusivity, \\begin{equation} \\eta_P = \\eta_A + \\eta, \\end{equation} which we parameterise by the constant nondimensional Pedersen diffusion parameter, $\\tilde{\\eta}_P$ (described in Appendix A).  Similarly, we define a nondimensional Hall diffusion parameter $\\tilde{\\eta}_H$ such that the Hall coefficient $\\eta_H$ scales with the surface density and thickness of the thin disc in a similar way to the Pedersen coefficient (see Appendix A). The direction of Hall diffusion depends upon the product $\\eta_H(\\mathbf{B}\\cdot\\Omega)$, where the sign of $\\eta_H$ depends on microphysics within the disc \\citep{wn1999}; in our simulations we vary the sign of $\\tilde{\\eta}_H$ in order to examine the effect of reversing the magnetic field. By studying the range of parameters that give stable accretion flow solutions it may be possible to place constraints on their physical values in observed accretion systems.  We look for solutions to the equations that take the form of power laws with respect to a similarity variable $x = r/c_st$ in the limit that $x\\to0$. These may then be thought of as valid in the limits $r\\to0$ or $t\\to\\infty$, that is, the innermost regions or late stages of the accretion flow. Only two physical solutions exist to these equations (see Appendix \\ref{Assim}): a secularly-evolving Keplerian disc and a magnetically-diluted near-free fall collapse. We present these in the following two sections. ",
        "conclusions": "We have presented two distinct power law similarity solutions to the MHD equations describing a flattened accretion flow with magnetic diffusion. The first of these represented a rotationally-supported disc through which gas is slowly accreted while the magnetic field diffuses outwards against the flow. The second was a free fall collapse onto the star in which almost all of the angular momentum has been removed from the fluid and no rotationally-supported disc may form. These have been used to show that even a small amount of Hall diffusion can dramatically change the dynamics of gravitational collapse."
    },
    "1208/1208.2834_arXiv.txt": {
        "abstract": "{% We report simultaneous multicolour observations in  5 bands $(UBVRI)$ of the  flickering variability of  the cataclysmic variable AE Aqr. Our aim is to estimate the parameters (colours, temperature, size) of the fireballs that produce the optical flares.  \\\\ The observed rise time of the optical flares is in the interval 220 - 440 sec. We estimate  the dereddened colours of the fireballs: $(U-B)_0$ in the range 0.8-1.4, $(B-V)_0 \\; \\sim$~0.03-0.24, $(V-I)_0 \\; \\sim$~0.26-0.78. We find for the fireballs  a temperature  in the range 10000 - 25000 K, mass  (7-90)$\\times 10^{19}$~g, size (3-7)$\\times 10^9$~cm  (using a distance of $d=86$~pc). These values refer to the peak of the flares observed in $UBVRI$ bands. \\\\ The data are available upon request from the authors. } ",
        "introduction": "AE Aquarii is an 11-12 mag cataclysmic variable, which was discovered in the optical by Zinner (1938) and was first associated with the DQ~Her (Intermediate Polar) stars by Patterson~(1979).  In  the AE~Aqr system a K0-K4 IV/V star transfers material through the L$_1$ nozzle toward  a magnetic white dwarf. The high-dispersion time-resolved absorption line spectro- scopy by  Echevarr{\\'{\\i}}a et al. (2008) gives the binary parameters as: white dwarf mass  $M_{WD} = 0.63 \\pm 0.05$ \\msun, secondary mass $M_2 = 0.37 \\pm 0.04 $ \\msun, binary separation $a = 2.33 \\pm 0.02$ R$_\\odot$, inclination $i \\approx 70^o$.   The light curve of AE~Aqr exhibits large flares and coherent oscillations of about 16 and 33~s in the optical and X-ray (Patterson 1979). It also exhibits radio and millimeter synchrotron emission (e.g. Bookbinder \\& Lamb 1987; Bastian, Dulk \\& Chanmugam~1988), and possibly even TeV $\\gamma$-rays (Bowden et al. 1992; Meintjes et al. 1992).  It has a long orbital period of 9.88~h (e.g. Casares et al. 1996) and a very short spin period of the white dwarf. On the basis of IUE spectra, Jameson, King \\& Sherrington (1980) reveal the most extreme C$_{IV}$/N$_{V}$ ratio of all Cataclysmic Variable stars, probably indicating strong carbon depletion and thus CNO cycling (Mauche, Lee \\& Kallman 1997). To arrive at such a state,  AE~Aqr is believed to be a former supersoft X-ray binary, in which the mass transfer rate in the recent past ($10^7$~yr) has been much higher than its current value (Schenker et al. 2002).   \\begin{table*} % \\footnotesize \\caption{Journal of observations. In the table are given as follows: the telescope, band, UT-start and UT-end of the run, exposure time, number of CCD images obtained, average magnitude in the corresponding band, minimum -- maximum magnitudes in each band, standard deviation of the mean, typical observational error.} \\begin{center} \\begin{tabular}{llrrrrcrrccccrlcr} \\hline  Telescope& Band       &   UT        & Exp-time  &  N$_{pts}$  &  average & min-max     & stdev  & err   \\\\ &               & start-end   & [sec]     &             & [mag]    & [mag]-[mag] & [mag]  & [mag] \\\\ \\hline {\\bf 2010 Aug 12 }  &     & JD2455421  \\\\  50/70 cm Schmidt    & $B$ & 23:20 - 00:32 & 60,90,120 & 42 & 12.338 & 12.214 - 12.397 & 0.047 & 0.010 \\\\ 60 cm Belogradchick  & $R$ & 23:00 - 00:37 & 40        & 81 & 10.717 & 10.607 - 10.754 & 0.027 & 0.007 \\\\ \\hline {\\bf 2010 Aug 13 } &     & JD2455422  \\\\ 2.0 m RCC         & $U$ & 22:23 - 00:28 & 180       & 19 & 11.577 & 10.985 - 12.335 & 0.336 & 0.025 \\\\ 50/70 cm Schmidt  & $B$ & 21:46 - 00:28 &  60       &137 & 12.226 & 11.694 - 12.430 & 0.169 & 0.008 \\\\ 2.0 m RCC         & $V$ & 22:23 - 00:28 &  10       &463 & 11.377 & 10.990 - 11.486 & 0.107 & 0.007 \\\\ 60 cm Belogradchick    & $R$ & 22:00 - 00:28 & 60,40     &181 & 10.699 & 10.422 - 10.797 & 0.077 & 0.006 \\\\ 60 cm Rozhen      & $I$ & 22:41 - 00:29 & 15        &388 & 10.177 &  9.987 - 10.261 & 0.052 & 0.007 \\\\ \\hline {\\bf 2010 Aug 14 } &     & JD2455423  \\\\ 2.0 m RCC         & $U$ &  18:57 - 23:43 & 120,180    & 67   & 12.073 & 10.824 - 12.942 & 0.527 & 0.022 & \\\\ 50/70 cm Schmidt  & $B$ &  18:44 - 23:44 & 60,120,180 & 245  & 12.186 & 11.672 - 12.450 & 0.144 & 0.011 & \\\\ 2.0 m RCC         & $V$ &  18:48 - 23:45 & 10         & 1106 & 11.346 & 11.056 - 11.831 & 0.080 & 0.007 & \\\\ 60 cm Belogradchick      & $R$ &  19:32 - 23:46 & 40,60      & 331  & 10.639 & 10.432 - 10.749 & 0.057 & 0.005 & \\\\ 60 cm  Rozhen     & $I$ &  18:24 - 23:42 & 15         & 1132 & 10.143 & 9.976 - 10.284  & 0.060 & 0.009 & \\\\ \\hline {\\bf 2010 Aug 16 } &     & JD2455425  \\\\ 60 cm  Rozhen      & $U$ &  18:31 - 19:59 & 120        & 24   & 12.510 & 12.013 - 13.100 & 0.278 & 0.085 & \\\\ 60 cm  Rozhen      & $B$ &  18:33 - 00:26 & 60,90      & 76   & 12.317 & 12.016 - 12.565 & 0.142 & 0.012 & \\\\ 60 cm  Belogradchick       & $V$ &  18:54 - 00:27 & 60,180\t & 68\t& 11.383 & 11.207 - 11.563 & 0.119 & 0.008 & \\\\ 60 cm  Belogradchick       & $R$ &  18:57 - 00:28 & 40\t & 71\t& 10.738 & 10.565 - 10.917 & 0.109 & 0.006 & \\\\ 60 cm  Belogradchick       & $I$ &  18:53 - 00:29 & 40\t & 70\t& 10.190 & 10.036 - 10.351 & 0.095 & 0.006 & \\\\ \\hline {\\bf 2011 Aug 31 } &     & JD2455805 \\\\ 2.0 m RCC          & $U$ &  20:59 - 23:46 & 90,120 & 72   & 11.368 &  9.841 - 12.476 & 0.685 & 0.039 & \\\\ 50/70 cm Schmidt   & $B$ &  20:54 - 23:40 & 30     & 270  & 11.777 & 10.767 - 12.339 & 0.376 & 0.032 & \\\\ 2.0 m RCC          & $V$ &  20:58 - 23:18 & 5      & 1090 & 11.079 & 10.355 - 11.738 & 0.228 & 0.012 & \\\\ 60 cm  Rozhen      & $R$ &  20:45 - 23:44 & 10     & 328  & 10.421 &  9.876 - 10.730 & 0.175 & 0.015 & \\\\ 60 cm  Rozhen      & $I$ &  20:45 - 23:44 & 10     & 327  &  9.950 &  9.546 - 10.195  & 0.129 & 0.015 & \\\\ \\hline  \\end{tabular} \\end{center} \\label{tab.journal} \\end{table*}   The Balmer emission lines vary both in strength and shape and they are not good tracers of the orbital motion of the white dwarf. This has led to the proposal of the magnetic propeller model (Wynn, King \\& Horne 1997). The gas stream emerging from the K2V companion star through the inner Lagrangian point L1 encounters a rapidly spinning magnetosphere. The rotating white dwarf in AE~Aqr ejects most of the matter transferred from the secondary in the form of  blobs (\\textquoteleft  fireballs\\textquoteright ). Only a small fraction ($\\sim$3\\%) of the mass flow at the magnetospheric radius eventually accretes on to the surface of the white dwarf, emphasizing the  effective magnetospheric propeller process in the system (Oruru \\& Meintjes 2012).  The white dwarf in AE Aqr is a fast rotator  having a spin period $P=33.08$~s, and  spinning down at a rate $5.64 \\times 10^{-14}$~s~s$^{-1}$ (de Jager et al. 1994; Mauche et al. 2011), corresponding to a spin-down luminosity of $6 \\times 10^{33}$~erg~s$^{-1}$  (Oruru \\& Meintjes 2011). A part of the spin-down power is consumed to expel the blobs.   In this paper we present simultaneous multicolour observations in  5 bands $(UBVRI)$ of the flares of AE~Aqr. The source of the flares is assumed to be the fireballs(blobs). We measure the fluxes emitted by the fireballs, their colours, and  calculate their mass, size, temperature and expansion rate. ",
        "conclusions": "Using 4 telescopes, we performed simultaneous  observations in  5 bands $(UBVRI)$ of the flare activity  of  the cataclysmic variable AE Aqr. The main results of our work are:  (1) For a total of  18.7 hours of observations we detected 14 flares. The rise time of the flares is in the range  220-440~s.  (2) For 5 individual fireballs, we calculated the peak fluxes in $(UBVRI)$   bands, rise times,  colours, and modeled their evolution. At the peaks of the optical emission, the dereddened colours of the fireballs are: $(U-B)_0$ in the range 0.8-1.4, $(B-V)_0 \\sim $    0.03-0.24.  (3) Adopting  the model of an isothermically expanding  ball of gas, we find for the individual fireballs a temperature in the range 10000 -- 25000~K, mass  7-90$\\times 10^{19}$~g,  size 3-7$\\times 10^9$~cm, constant of expansion  1/H = 210 - 370~s$^{-1}$ (using a distance to AE~Aqr of $d=86$~pc and  interstellar extinction $E_{B-V} =0.05$). These values refer to the peak of the flares observed in the optical bands.  We also briefly discuss  the possible relation of the blobs of AE~Aqr with those detected in two symbiotic stars."
    },
    "1208/1208.2233_arXiv.txt": {
        "abstract": "We present the design and performance of the multi-object fiber spectrographs for the Sloan Digital Sky Survey (SDSS) and their upgrade for the Baryon Oscillation Spectroscopic Survey (BOSS). Originally commissioned in Fall 1999 on the 2.5-m aperture Sloan Telescope at Apache Point Observatory, the spectrographs produced more than 1.5 million spectra for the SDSS and SDSS-II surveys, enabling a wide variety of Galactic and extra-galactic science including the first observation of baryon acoustic oscillations in 2005.  The spectrographs were upgraded in 2009 and are currently in use for BOSS, the flagship survey of the third-generation SDSS-III project. BOSS will measure redshifts of 1.35 million massive galaxies to redshift 0.7 and Lyman-$\\alpha$ absorption of 160,000 high redshift quasars over 10,000 square degrees of sky, making percent level measurements of the absolute cosmic distance scale of the Universe and placing tight constraints on the equation of state of dark energy.  The twin multi-object fiber spectrographs utilize a simple optical layout with reflective collimators, gratings, all-refractive cameras, and state-of-the-art CCD detectors to produce hundreds of spectra simultaneously in two channels over a bandpass covering the near ultraviolet to the near infrared, with a resolving power $R = \\lambda$/FWHM $\\sim 2000$. Building on proven heritage, the spectrographs were upgraded for BOSS with volume-phase holographic gratings and modern CCD detectors, improving the peak throughput by nearly a factor of two, extending the bandpass to cover 360 $<$ $\\lambda$ $<$ 1000 nm, and  increasing the number of fibers from 640 to 1000 per exposure. In this paper we describe the original SDSS spectrograph design and the upgrades implemented for BOSS, and document the predicted and measured performances. ",
        "introduction": "\\label{sec:intro}    The Sloan Digital Sky Survey \\citep[SDSS;][]{york00a} project was conceived in the mid-1980s as an ambitious endeavor to understand the large-scale structure of the Universe. SDSS and its extension, SDSS-II, conducted a coordinated imaging and spectroscopic survey from 2000-2008 over approximately 10,000 deg$^{2}$ of high Galactic latitude sky. Now in its third phase of operation, SDSS is one of the most successful projects in the history of astronomy. The survey has produced an enormous catalog consisting of five-band digital images that include nearly one billion unique objects, and spectra of 930,000 galaxies, 120,000 quasars, and 460,000 stars, all publicly available \\citep[][and references therein]{abazajian09a}.  To obtain these imaging and spectroscopic data, a dedicated 2.5m telescope \\citep{gunn06a}, wide-field mosaic CCD camera \\citep{gunn98a}, and twin multi-object fiber spectrographs were constructed and installed at the Apache Point Observatory (APO) in Sunspot, New Mexico. The telescope, built to accommodate the requirements for both imaging and spectroscopy, is shared by the camera and spectrographs, which mount at the Cassegrain focus. The imaging survey was carried out on clear, dark nights with good seeing using the 120 mega-pixel camera, which operated in drift-scanning mode using a $5 \\times 6$ array of $2048 \\times 2048$ pixel detectors to obtain $ugriz$ \\citep{fukugita96a}, photometry. The imaging data, once reduced and calibrated \\citep{smith02a,pier03a,ivezic04a,tucker06a,padmanabhan08a}, were used for spectroscopic target selection. Spectroscopy was performed using the two multi-object fiber spectrographs, collecting 640 spectra over the 3$^\\circ$ diameter field in one exposure.  In this paper, we describe the design and performance of the SDSS spectrographs, and their recent upgrade for the Baryon Oscillation Spectroscopic Survey \\citep[BOSS][]{schlegel09a,dawson13a}. BOSS is the flagship survey in the third-generation SDSS-III program currently underway at the 2.5m SDSS telescope \\citep{eisenstein11a}. BOSS will measure the cosmic expansion history of the universe to percent-level precision by mapping an immense volume of sky to obtain the spatial distributions of galaxies and quasars, and from it, the characteristic scale imprinted by baryon acoustic oscillations (BAO) in the early universe \\citep[for a review of BAO with a respect to other cosmological probes, see][]{weinberg12a}. A measure of the scale at low redshifts, out to $z \\sim 0.7$, will be obtained by carrying out a redshift survey of 1.35 million massive galaxies from 10,000 deg$^{2}$ of SDSS data. BOSS will also observe Lyman-$\\alpha$ absorption in the spectra of 160,000 high-redshift quasars to measure large scale structure at redshifts of $z \\sim 2.5$.  Each SDSS spectrograph utilizes a dual-channel design with a common reflecting collimator and a dichroic to split the beam into a blue channel and a red channel. In each channel, just downstream of the dichroic, a transmitting grism disperses the light, which is imaged by an all-refractive camera onto a CCD. For BOSS the basic optical design has been retained, with several improvements. The ruled gratings have been replaced by volume-phase holographic (VPH) grisms (gratings sandwiched between two prisms) and the CCDs have been replaced with more modern devices. These changes produce a significant improvement in throughput and a modest extension of the wavelength range in both the blue and red channels. Additionally, smaller diameter fibers that are better matched to the angular scale of BOSS targets have been installed, allowing the total number of simultaneous spectra obtained from the two spectrographs to be increased from 640 in the original design to 1000 in the BOSS configuration.   The remainder of this paper is organized as follows. In Section 2 we begin by describing the design and construction of the original SDSS spectrographs in some detail, published here for the first time. This is followed in Section 3 by a discussion of the spectrograph upgrades completed in 2009 for BOSS. The performance of both the original SDSS spectrographs and the upgraded BOSS design is presented in Section 4. Finally, some highlights of the scientific research enabled by these instruments is provided in Section 5. ",
        "conclusions": "\\label{sec:conclusion}   The SDSS spectrographs were designed to be high-throughput, robust instruments, to conduct an extragalactic survey producing a million redshifts over a five year lifetime. These goals have clearly been met, with 1.6 million spectra of stars, galaxies, and quasars over 2880 plates, primarily in the northern Galactic cap at high Galactic latitudes. Spectra over 9274 deg$^2$ were collected in nine years of operation between the start of SDSS-I in 1999 and the completion of the first of two extended phases, SDSS-II, including SEGUE which produced roughly 250,000 stellar spectra between 2004-2008 \\citep{yanny09a}. These data were released to the astronomy community nearly every year; the full data set from SDSS-I and II is included in DR8 \\citep{aihara11a}. Aside from a few minor technical issues where the instruments did not meet predicted design performance, the spectrographs proved to be quite reliable and exceptionally productive instruments. With improvements in CCD technology, the use of VPH gratings, and enhancements to the optics (i.e. fluid coupled camera triplets, new dichroics, and higher reflectivity collimator coatings) the instrumental throughput was enhanced significantly, allowing the use of smaller diameter fibers and making the BOSS survey possible. Along with improvements to the guide camera, modifications to improve flexure, and a host of minor upgrades, the BOSS spectrographs have proven to be as productive as the original SDSS design and have many years of life remaining.  The spectroscopic data from SDSS have enabled studies across a broad range of astronomical disciplines including the evolution and clustering of galaxies \\citep[e.g.,][]{kauffmann04a,tegmark04a}, gravitational lensing \\citep[e.g.,][]{bolton06a}, the properties of quasars \\citep[e.g.,][]{vandenberk01a}, and stellar astrophysics \\citep[e.g.,][]{west08a}. One of the prominent scientific contributions from SDSS and SDSS-II data is the discovery of acoustic oscillation signatures in the clustering of galaxies \\citep{eisenstein05a}, opening the door to a new method of cosmological measurement. SDSS imaging and spectroscopic data have been included in more than 3500 refereed papers.  As described in detail in \\citet{dawson13a}, the BOSS survey will cover 10,000 deg$^{2}$ in five years, obtaining spectra for 1.35 million luminous galaxies with redshifts $0.15<z<0.7$, and approximately 160,000 quasars with redshifts between $2.15<z<3.5$. The galaxies provide a highly biased tracer of matter, and at these densities over this area provide percent-level distance scale measurements through studies of BAO. The quasars provide a backlight to illuminate neutral hydrogen through Lyman-$\\alpha$ absorption along the line of sight, thereby mapping large-scale structure in the foreground of each quasar. These distance measurements will represent a significant improvement on the accuracy of existing BAO measurements, and will tighten the constraints on the dark energy equation of state, $w=p/\\rho$, where $p$ is pressure and $\\rho$ is density, and the evolution of $w$ with time.  Detection of the BAO feature in quasar spectra would allow BOSS to make the first characterization of dark energy at early times, when dark energy should be sub-dominant according to the prevailing $\\Lambda$CDM theory.  BOSS completed imaging of the SGC spectroscopic footprint in 2008 and 2009; spectroscopic observations began December 5, 2009 after a three month commissioning phase. The algorithm for selecting galaxy targets is outlined briefly in \\citet{white11a} and \\citet{eisenstein11a}. The quasar target selection is described in \\citet{ross12a} using photometry from the SDSS imaging survey supplemented with data from programs at ultraviolet, infrared, and radio wavelengths to enhance the detection efficiency. These first two years of BOSS spectra include 324,198 unique galaxy targets ($z>0.43$) with a 98.7\\% rate of successful classification, and 103,729 unique lower redshift galaxy targets ($0.15<z<0.43$) with a 99.9\\% successful rate of classification. There were more than 85,000 quasars observed, 61,933 of which were at $z>2.15$. The spectra of these objects are classified as described in \\citet{bolton12a} and each quasar is visually inspected and recorded as explained in \\citet{paris12a}. Data from the first two years of the survey were included in DR9 \\citep{ahn12a} and include the spectra from 831 plates.  The higher redshift sample of galaxies from DR9 has already been used to fit the position of the acoustic peak at an effective redshift $z=0.57$ \\citep{anderson12a}. With 1.7\\% accuracy, this is the most precise distance constraint ever obtained from a galaxy survey at any redshift. The potential for distance measurements at $z>2$ using Lyman-$\\alpha$ forest absorption from quasars is promising; the first detection of flux correlations across widely separated sightlines was obtained in the first year of BOSS data to comoving separations of 60 $h^{-1}$Mpc \\citep{slosar11a}. More recently, the acoustic peak was measured in the Lyman-$\\alpha$ forest, two complimentary analyses produced distance measurements at $z>2$ with 3\\% \\citep{busca13a} and 2\\% \\citep{slosar13a} precision. BOSS is on pace to complete the 10,000 deg$^2$ spectroscopic footprint by the middle of 2014, providing new cosmology constraints and a wealth of data for studies of galaxy evolution and quasar physics.  The optical spectrographs at APO will have unique capabilities for widefield, multiplexed spectroscopy well beyond the 2014 completion of the BOSS survey. The SDSS collaboration has proposed an ``After SDSS III'' (AS3) program for 2014-2020 that will capitalize on this resource. Four separate programs in AS3 will use the BOSS spectrograph, likely with only minor modifications. The extended Baryon Oscillation Spectroscopic Survey (eBOSS) will probe dark energy and fundamental physics by making distance measurements with BAO in the redshift range $0.6<z<2$. The Time-Domain Spectroscopic Survey (TDSS) will obtain spectra of 50,000 -- 100,000 Galactic and extragalactic variable sources with simultaneous, multi-epoch optical imaging data. In a program to follow up x-ray sources in new 0.2\u20138 keV data obtained from the extended ROentgen Survey with an Imaging Telescope Array \\citep[eROSITA:][]{predehl10a}, the SPectroscopic IDentification of eROSITA Sources (SPIDERS) survey will follow up 50,000 -- 100,000 objects. Finally, Mapping Nearby Galaxies at Apache Point Observatory (MaNGA) will perform spatially resolved spectroscopy on approximately 10,000 nearby galaxies using 15 integral field units integrated into new BOSS-like cartridges. As with the original SDSS spectroscopic survey, these four surveys will provide a premier data sample for astrophysical studies from Galactic to cosmological scales."
    },
    "1208/1208.3003_arXiv.txt": {
        "abstract": "We present a homogeneous chemical abundance analysis of 16 elements in 190 metal-poor Galactic halo stars (38 program and 152 literature objects). The sample includes 171 stars with [Fe/H] $\\le$ $-$2.5, of which 86 are extremely metal poor, [Fe/H] $\\le$ $-$3.0. Our program stars include ten new objects with [Fe/H] $\\le$ $-$3.5.  We identify a sample of ``normal'' metal-poor stars and measure the trends between [X/Fe] and [Fe/H], as well as the dispersion about the mean trend for this sample. Using this mean trend, we identify objects that are chemically peculiar relative to ``normal'' stars at the same metallicity. These chemically unusual stars include CEMP-no objects, one star with high [Si/Fe], another with high [Ba/Sr], and one with unusually low [X/Fe] for all elements heavier than Na.  The Sr and Ba abundances indicate that there may be two nucleosynthetic processes at lowest metallicity that are distinct from the main $r$-process.  Finally, for many elements, we find a significant trend between [X/Fe] versus \\teff\\ which likely reflects non-LTE and/or 3D effects. Such trends demonstrate that care must be exercised when using abundance measurements in metal-poor stars to constrain chemical evolution and/or nucleosynthesis predictions. ",
        "introduction": "\\label{sec:intro}  The atmospheres of low-mass stars contain detailed information on the chemical composition of the interstellar medium at the time and place of their birth. Thus, studies of the most metal-poor stars of the Galactic halo arguably offer the best means with which to understand the properties of the first stars (e.g., \\citealt{chamberlain51,wallerstein63,carney81,bessell84,mcwilliam95,ryan96,norris01,johnson02,cayrel04,beers05,cohen08,lai08,frebel11}). In recent times, four stars with an iron content less than $\\sim1/30,000$ that of the Sun have been discovered -- HE~0107$-$5240 \\citep{christlieb02,christlieb04}, HE~1327$-$2326 \\citep{frebel05,aoki06}, and HE~0557$-$4840 \\citep{norris07}, within the Hamburg/ESO Survey (HES; \\citealt{hes}), and SDSS J102915+172927 \\citep{caffau11}, in the Sloan Digital Sky Survey Data Release 7 \\citep{sdss7}.  The abundance patterns of these stars can constrain yields from zero metallicity supernovae (e.g., \\citealt{limongi03,umeda05,meynet06,tominaga07,heger10,limongi12}).  Chemical abundance studies of metal-poor stars with increasing accuracy and precision have, in some cases, revealed extremely small scatter in abundance ratios at low metallicity \\citep{cayrel04,arnone05}. Such results place strong constraints on the yields of the progenitor stars, as well as on the relative contributions of intrinsic scatter and measurement errors. Meanwhile, parallel studies have started to identify a variety of chemically diverse objects (e.g., \\citealt{aoki05,cohen07,lai09}), which may indicate that ``we are beginning to see the anticipated and long sought stochastic effects of individual supernova events contributing to the Fe-peak material within a single star'' \\citep{cohen08}.  As the numbers of metal-poor stars with detailed chemical abundance measurements have grown, databases have been compiled \\citep{saga,saga2,frebel10} which facilitate studies of the global characteristics of metal-poor stars (see also \\citealt{roederer09}, who assembled and studied a compilation of nearly 700 halo stars). However, it is difficult, perhaps impossible, to understand whether the trends and dispersions in [X/Fe]\\footnote{We adopt the usual spectroscopic notations that [A/B]~$\\equiv$ log$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\star}$~-- log$_{\\rm 10}$(N$_{\\rm A}$/N$_{\\rm B}$)$_{\\sun}$, and that log~$\\varepsilon$(B)~=~A(B)~$\\equiv$ log$_{\\rm 10}$(N$_{\\rm B}$/N$_{\\rm H}$)~$+$~12.00, for elements A and B.} versus [Fe/H] are real, or an artefact of an inhomogeneous comparison. Moreover, in the absence of a careful, homogeneous analysis, subtle effects may be overlooked.   In order to make further progress in this field, there is a clear need for a detailed homogeneous chemical abundance analysis of a large sample of metal-poor stars. To this end, \\citet{barklem05} studied 253 stars, presenting chemical abundance measurements for some 22 elements. Their study included 49 new stars with [Fe/H] $<$ $-$3, but only one of which had [Fe/H] $\\le$ $-$3.5, thereby highlighting the difficulty of finding stars in this metallicity regime.  This is the second paper in our series on the discovery and analysis of the most metal-poor stars. Here, we present a homogeneous chemical abundance analysis for 38 program stars and a further 152 literature stars. Our program stars include ten new objects with [Fe/H] $\\le$ $-$3.5, and the combined sample includes 86 extremely metal-poor stars, [Fe/H] $\\le$ $-$3.0. To our knowledge, this represents one of the largest homogeneous chemical abundance analyses of the most metal-poor stars to date based on model atmosphere analysis of equivalent widths measured in high-resolution, high signal-to-noise ratio (S/N) spectra. The outline of the paper is as follows. In Section~\\ref{sec:obs}, we describe the analysis of the 38 program stars. In Section~\\ref{sec:firststars}, we compare our chemical abundances with those of the {\\sc First Stars} group \\citep{cayrel04,spite05,francois07}. In Section~\\ref{sec:lit}, we describe our homogeneous re-analysis of 207 literature stars; in Section~\\ref{sec:comp}, we compare these results with the literature values. In Section~\\ref{sec:nlte}, we consider non-LTE effects. Finally, our results, interpretation, and conclusions are presented in Sections~\\ref{sec:results} and \\ref{sec:conc}. ",
        "conclusions": ""
    },
    "1208/1208.3529_arXiv.txt": {
        "abstract": "The physical properties inferred from the spectral energy distributions of $z>3$ galaxies have been influential in shaping our understanding of early galaxy formation and the role galaxies may play in cosmic reionization.  Of particular importance is the stellar mass density at early times which represents the integral of earlier star formation. An important puzzle arising from the measurements so far reported is that the specific star formation rates (sSFR) evolve far less rapidly than expected in most theoretical models. Yet the observations underpinning these results remain very uncertain, owing in part to the possible contamination of rest-optical broadband light from strong nebular emission lines. To quantify the contribution of nebular emission to broad-band fluxes, we investigate the spectral energy distributions of 92 spectroscopically-confirmed galaxies in the redshift range $3.8<z<5.0$ chosen because the H$\\alpha$ line lies within the {\\it Spitzer}/IRAC 3.6$\\mu$m filter. We demonstrate that the 3.6$\\mu$m flux is systematically in excess of that expected from stellar continuum alone, which we derive by fitting the SED with population synthesis models. No such excess is seen in a control sample of spectroscopically-confirmed galaxies with $3.1<z<3.6$ in which there is no nebular contamination in the IRAC filters. From the distribution of our 3.6$\\mu$m flux excesses, we derive an H$\\alpha$ equivalent width distribution and consider the implications both for the derived stellar masses and the sSFR evolution. The mean rest-frame H$\\alpha$ equivalent width we infer at $3.8<z<5.0$ (270~\\AA) indicates that nebular emission contributes at least 30\\% of the 3.6$\\mu$m flux and, by implication, nebular emission is likely to have a much greater impact for galaxies with $z\\simeq 6-7$ where both warm IRAC filters are contaminated. Via our empirically-derived  equivalent width distribution we correct the available stellar mass densities and show that the sSFR evolves more rapidly at $z>4$ than previously thought, supporting up to a 5$\\times$ increase between $z\\simeq 2$ and 7.  Such a trend is much closer to theoretical expectations.  Given our findings, we discuss the prospects for verifying quantitatively the nebular emission line strengths prior to the launch of the James Webb Space Telescope. ",
        "introduction": "\\label{sec:intro}  Through detailed photometry of Lyman break galaxies (LBGs) undertaken with the Hubble and Spitzer Space Telescopes, much has been learned regarding the physical properties of galaxies beyond a redshift $z\\simeq$3.   Stellar masses and star formation rates have now been inferred from broadband photometric spectral energy distributions (SEDs) for thousands of galaxies spanning the redshift range $3<z<7$ (e.g., \\citealt{Egami05,Eyles05,Labbe06,Eyles07,Stark07,Stark09,Ono10, Gonzalez10,Labbe10a,Labbe10b,Gonzalez11a,Lee12,Reddy12,Curtis-Lake12}).  The stellar mass density derived from these studies has proven a useful integrated constraint on the contribution of galaxies to reionization (e.g., \\citealt{Robertson10}), while the evolution of physical properties has provided insight into the processes which govern the assembly of early galaxies (e.g., \\citealt{Finlator11,Dave11,Dave12}).  A potentially significant puzzle has recently emerged from these studies through measurement of the specific star formation rate (sSFR) at $z>2$.  Current observations demonstrate that between $z\\simeq 2$ and $z\\simeq 7$, the sSFR in galaxies of fixed stellar mass does not evolve strongly (e.g., \\,\\citealt{Stark09,Gonzalez10}), with recent estimates indicating at most a factor of two increase between $z\\simeq 2$ and 7 (e.g., \\citealt{Bouwens12b,Reddy12}).   This is in contrast to simple expectations from semi-analytic models and numerical simulations (e.g., \\citealt{Weinmann11,Dave11,Dave12}) which predict that the sSFR should closely match the inflow rate of baryonic material. As this mass inflow rate is thought to increase with redshift as $\\dot{M}/M\\simeq (1+z)^{2.25}$ \\citep{Neistein08,Dekel09}, we should expect nearly a 10$\\times$ increase in sSFR in galaxies of fixed stellar mass over $2<z<7$, in marked contrast to the observations.  The physical cause of the discrepancy associated with the sSFR evolution remains unclear. As discussed previously (e.g., \\citealt{Bouche10,Dutton10,Weinmann11,Dave11,Dave12,Reddy12}), a plateau in the redshift dependence of the sSFR would suggest star formation is more inefficient at $z>6$ than at $z\\simeq 2$.  Various physical processes  might be invoked to impede star formation, such as the inefficient formation of molecular hydrogen in low metallicity galaxies (e.g., \\citealt{Robertson08,Gnedin09,Krumholz12}) or an increase in the efficiency with which  cold gas is removed via large-scale outflows.  Irrespective of any mechanism that might inhibit star formation at early times, it is difficult to reconcile such an inefficiency with the notion that early galaxies provide the ionizing photons responsible for reionization \\citep{Robertson10}. For example, the steep faint end slope of the ultraviolet luminosity function (UV LF) at $z>6$ \\citep{Bouwens11} implies that star formation in low mass dark matter halos becomes {\\it more efficient} at earlier times  (e.g., \\citealt{Trenti10}), in contrast to the implications of the sSFR measurements.  Given these difficulties, it is prudent that we reconsider the accuracy of the data that is used to infer the sSFR and its evolution.  The two basic ingredients are the star formation rates and the stellar masses. The $z>4$ measurements have indeed changed since the original articles (e.g., \\citealt{Stark09,Gonzalez10}), mostly as a result of improved dust corrections following improved near-infrared photometry (e.g., \\citealt{Bouwens12b}). The new dust corrections have served to increase the $z\\simeq4$ sSFR measurements by a factor $\\simeq$ 2. However, since negligible extinction is inferred at $z\\simeq 6-7$, the sSFR still remains constant over $4<z<7$ \\citep{Bouwens12b} although a factor of 2 higher than at $z\\simeq 2-3$ \\citep{Reddy12}.  A potentially more important problem is the possible contribution of rest-frame optical nebular emission lines (e.g., [O II], [O III], H$\\alpha$) to the broad-band fluxes used to infer the stellar masses. Such emission lines, could significantly affect the inferred amplitude of a Balmer Break, leading to an overestimate of the stellar mass and thereby an underestimate of the sSFR. Figure 1 illustrates how the various nebular emission lines contaminate the key photometric filters as a function of redshift.  Beyond $z\\simeq4$, the key filters of interest in the determination of stellar masses are the {\\it Spitzer}/IRAC warm bands at 3.6$\\mu$m and 4.5$\\mu$m. It is particularly striking that, at $z\\gsim 5$, the strongest rest-frame optical nebular lines ([OIII]$\\lambda5007$ and H$\\alpha$) contaminate both {\\it Spitzer}/IRAC filters. Although many $z\\gsim 5$ galaxies are detected with {\\it Spitzer} (e.g., \\citealt{Egami05,Eyles05,Stark09,Labbe10a,Labbe10b,Gonzalez10,Gonzalez11a,Richard11b}), contamination by nebular emission could significantly affect the interpretation of their SEDs.  Accounting for nebular emission in the SEDs of high redshift galaxies has been considered by several earlier works (e.g., \\citealt{Schaerer09,Schaerer10,Ono10,deBarros12}).  In general terms, their approach has been to use `forward modeling' techniques based on adding nebular emission contributions to stellar population synthesis models in order to demonstrate the possible implications of its inclusion. However, such ``nebular+stellar\" model fits cannot provide a precise unambiguous measure of nebular contamination for several reasons. Firstly, there are numerous uncertainties in how the contribution of nebular emission should be added. These include the nebular extinction law and ionizing photon escape fraction. Secondly, for galaxies at $z>5$, for which there is no uncontaminated measure of the stellar continuum (Figure 1), the uncertainties are particularly large. Finally, and perhaps most importantly, without a spectroscopic redshift, addressing both the nebular contamination and the photometric redshift of the galaxy from the same photometric data leads to great uncertainties; there is no {\\it a priori} indication of which photometric bands are contaminated by nebular emission.  Fortunately, by virtue of our deep Keck spectroscopic survey (Stark et al. 2010, Stark et al. 2011, Jones et al. 2012, Schenker et al. 2012) and our nebular+stellar population synthesis code (Robertson et al. 2010), we can use the availability of HST-Spitzer SEDs to make progress in addressing this issue. While the question of contamination by nebular emission at $z>5$ must await the infrared spectroscopic capabilities of James Webb Space Telescope, we can test our spectroscopic range $3.8<z<5.0$ for contamination by H$\\alpha$ in the {\\it Spitzer}/IRAC 3.6$\\mu$m broadband filter. Our approach follows that of \\cite{Shim11} who demonstrated that galaxies in this redshift window are typically significantly brighter at 3.6$\\mu$m than at 4.5$\\mu$m.  By comparing their flux density at 3.6$\\mu$m to that expected from stellar continuum alone, \\cite{Shim11} argued that many galaxies at $z>4$ show evidence for  strong H$\\alpha$ emission, with typical EWs significantly greater than those seen at $z\\simeq 2$.  Here we seek to apply a similar technique to our spectroscopic sample \\citep{Stark10,Stark11} with the goal of estimating the {\\it distribution of H$\\alpha$ equivalent widths} present in galaxies at $3.8<z<5.0$.  Equipped with this external constraint on the strength of nebular emission, we can then determine how stellar masses and the sSFR of $z>4$ galaxies are likely to be altered by emission line contamination.   In particular, we will explore whether our estimated degree of nebular contamination could be sufficient at the highest redshifts to permit a rapid rise in the redshift-dependent sSFR as expected from theoretical models.  The present paper is organized as follows.   In \\S2, we discuss the selection of the spectroscopic sample used in our analysis. In \\S3, we introduce the details of our SED fitting procedure used to estimate the strength of nebular emission lines in the various filters.  In \\S4 we use our spectroscopic sample to estimate the equivalent width distribution of H$\\alpha$ in the redshift range $3.8<z<5.0$ and then use these measurements to assess the impact of nebular emission on the derived stellar masses and star formation rates of $z>4$ galaxies.  In \\S5, we discuss the impact that our findings have for the evolution in the integrated stellar mass density and specific star formation rates of galaxies at $z>4$.  Throughout this paper we adopt a $\\Lambda$-dominated, flat universe with $\\Omega_{\\Lambda}=0.7$, $\\Omega_{M}=0.3$ and $H_{0}=70\\,h_{70}  {\\rm km\\,s}^{-1}\\,{\\rm Mpc}^{-1}$. All magnitudes are quoted in the AB system \\citep{Oke83}.  \\begin{figure} \\epsscale{1.1} \\plotone{f1.eps} \\caption{Emission line contamination of broadband photometry.  Colored stripes denote redshift ranges over which emission lines contaminate the K$_s$-band (dark blue), IRAC 3.6 $\\mu$m (yellow), and IRAC 4.5 $\\mu$m (red).    H$\\alpha$ emission is expected in the 3.6$\\mu$m filter at $3.8<z<5.0$. Note that at $5\\lsim z\\lsim 7$, both IRAC filters used for measuring stellar masses are contaminated by emission lines.  Beyond $z\\simeq 7$, only the 4.5$\\mu$m filter is contaminated by strong nebular emission.   } \\end{figure} ",
        "conclusions": "\\label{sec:conclusion}  Measurement of the evolving stellar mass and sSFR distributions have proven critical to our understanding of early galaxy assembly and the UV photon budget of reionization-era galaxies.   Recently, it has become clear that many of these early estimates might be significantly in error due to the contamination of the {\\it Spitzer}/IRAC bandpasses by nebular emission lines.    Knowledge of the strength of these emission lines is necessary for robust determinations of the stellar mass density and sSFR evolution.  As these emission lines are shifted out of the observed atmospheric window, direct spectroscopic measurements will not be feasible until JWST.  In this paper, we present a method which enables constraints on nebular emission at $z>4$ by combining large spectroscopic samples and deep {\\it Spitzer} photometry.   Like Shim et al. (2011), we focus on the redshift range $3.8<z<5.0$, over which  the IRAC [3.6] filter is contaminated by strong emission lines (H$\\alpha$, [NII], [SII]) while the [4.5] filter is free of nebular contamination.  Examining a carefully-selected subset of 45 galaxies, we find that the 3.6$\\mu$m flux is systematically in excess of the expected stellar continuum flux, revealing the presence of strong nebular emission.  No excess is seen in a spectroscopic sample at $3.1<z<3.6$, a redshift range over which no strong emission lines contaminate the IRAC filters.   We use the photometric excesses in the contaminated [3.6] filter to estimate the equivalent width distribution of H$\\alpha$ emission at $3.8<z<5.0$.   Equipped with this measure of nebular emission at high-redshift, we re-evaluate the evolution in the sSFR and stellar mass density at $z\\gsim 4$.   Our primary conclusions from this analysis are summarized below.  1.  We find that the mean equivalent width of emission lines contaminating the [3.6] filter is $\\rm{<log_{10} (W_{3.6}/\\AA)>}$ $\\simeq 2.57-2.73$. We estimate that $\\simeq 76$\\% of this signal arises from H$\\alpha$, implying an average H$\\alpha$ equivalent width of $\\rm{<log_{10} (W_{H\\alpha}/\\AA)>}$ $\\simeq 2.45-2.61$ at $3.8<z<5.0$.  2.  The mean H$\\alpha$ equivalent width inferred at $3.8<z<5.0$ appears greater than that for similar star-forming samples at lower redshifts.   While definitive knowledge of the evolution in the H$\\alpha$ EW surely awaits direct spectroscopic measurement,   the evolution we infer over $2\\lsim z\\lsim 5$ is certainly consistent with the (1+z)$^{1.8}$ power law derived in Fumagalli et al. (2012).   This likely reflects an increase in the sSFR at $z\\gsim 2$, and importantly suggests that the EW of nebular emission continues to increase at $z\\gsim 5$.  3.  Using the H$\\alpha$ EW distribution we derive at $3.8<z<5.0$, we explore how nebular contamination is likely to affect the physical properties of galaxies at $z>3$. We find that the stellar masses are reduced, on average, by 1.1, 1.3, 1.6, and 2.4$\\times$ for dropout samples with mean redshifts of $z\\simeq \\rm{4,~5,~6,~and~7}$, respectively.   If the equivalent widths of nebular lines continue to increase in amplitude at $z\\gsim 5$, we estimate that the reduction in the stellar masses are likely to increase to 1.9 and 4.4$\\times$ at $z\\simeq 6$ and 7. We note that these corrections are representative for average measures of the dropout populations and not individual galaxies.  4.   As the stellar mass density provides a valuable integrated measure of early star formation, constraints on the level of nebular contamination are critical to our knowledge of the ionizing photon budget of galaxies throughout the reionization era.    After correcting for nebular emission contamination, we find a factor of $\\simeq 2\\times$ reduction from previous estimates.  The downward revisions to the stellar mass density improve consistency with expectations from the integrated star formation rate density.     Extending such nebular-corrected measurements to emerging galaxy samples at $z\\simeq 8-9$ will yield an integral constraint on the UV photon budget during $z\\simeq 10-15$.  5.    Whereas previous derivations showed little evolution in the sSFR of fixed mass galaxies over $2<z<7$,  we demonstrate that after accounting for nebular emission and correcting for dust, the sSFR increases by 4-6$\\times$ over $2<z<7$.    The absolute sSFR values inferred at $z\\gsim 5$ appear largely similar with predictions from simulations.   While there certainly remains room for improvement (in both the data and the modelling), the increase in the sSFR at $z\\gsim 4$ seems consistent with a picture whereby increasing baryon accretion rates at larger redshift translate into larger sSFR in galaxies of a fixed stellar mass.  \\subsection*"
    },
    "1208/1208.4286_arXiv.txt": {
        "abstract": "Determining the temperature distribution of coronal plasmas can provide stringent constraints on coronal heating. Current observations with the Extreme ultraviolet Imaging Spectrograph onboard Hinode and the Atmospheric Imaging Assembly onboard the Solar Dynamics Observatory provide diagnostics of the emission measure distribution (EMD) of the coronal plasma.  Here we test the reliability of temperature diagnostics using 3D radiative MHD simulations.   We produce synthetic observables from the models, and apply the Monte Carlo Markov chain EMD diagnostic. By comparing the derived EMDs with the ``true'' distributions from the model we assess the limitations of the diagnostics, as a function of the plasma parameters and of the signal-to-noise of the data.  We find that EMDs derived from EIS synthetic data reproduce some general characteristics of the true distributions, but usually show differences from the true EMDs that are much larger than the estimated uncertainties suggest, especially when structures with significantly different density overlap along the line-of-sight. When using AIA synthetic data the derived EMDs reproduce the true EMDs much less accurately, especially for broad EMDs. The differences between the two instruments are due to the: (1) smaller number of constraints provided by AIA data, (2) broad temperature response function of the AIA channels which provide looser constraints to the temperature distribution.  Our results suggest that EMDs derived from current observatories may often show significant discrepancies from the true EMDs, rendering their interpretation fraught with uncertainty. These inherent limitations to the method should be carefully considered when using these distributions to constrain coronal heating. ",
        "introduction": "\\label{s:intro}  The heating mechanism that is responsible for the million degree solar corona remains unknown, though several candidates exist. It is one of the most important open issues in astrophysics. Constraining the properties of this heating mechanism is a difficult task, but usually performed through spectral and imaging observations of the solar corona, which provide diagnostics of the plasma temperature distribution. The latter has important implications for the energy balance of the corona (see e.g., \\citealt{Klimchuk06,Reale10} and references therein).  The thermal distribution of the plasma - or emission measure distribution, EMD (in section \\ref{s:method} we define the relationship between EMD and the differential emission measure, DEM, which is also often used to describe the plasma thermal distribution) - is crucial to test heating models. For instance, the EMD in coronal loops, strongly depends on the spatial and temporal properties of the energy release \\citep[e.g.,][]{Klimchuk01,Cargill04,Testa05}. The presence of a high temperature ($T \\gtrsim 5$MK) component, e.g., in the EMD of an active region, is a good tracer of the properties of the heating \\citep[e.g.,][]{Patsourakos06}, and it has recently been addressed by several studies \\citep[e.g.,][]{Patsourakos06,Reale09,Reale09b,Ko09,Schmelz09a,Testa11,Testa12b}. Also, the thermal distribution in the cross-field direction can help discern between a monolithic, single strand, loop (isothermal at a given location along the loop; e.g., \\citealt{Aschwanden00,Aschwanden05,Aschwanden11, Delzanna03,Landi02,Landi06,Landi08}), and a loop structure composed of several strands (multi-thermal plasma; e.g., \\citealt{Schmelz01,Schmelz05,Warren09,Brooks09}).  \\begin{figure*}[!t]\\vspace{0.5cm} \\centerline{\\includegraphics[scale=0.45]{fig1a.ps} \\includegraphics[scale=0.45]{fig1b.ps}} \\caption{ Three-dimensional snapshot of one of the 3D rMHD models (model H): top view (left panel) and side view (right panel).   Some selected magnetic field lines are shown in red to give an indication of the magnetic field topology in the corona. The isosurfaces of \\fexii\\ and \\caxvii\\ emission are shown in green and blue respectively. We also show at the photospheric layer (z=0~Mm) the strength of the vertical magnetic field in blue-red color scale. \\label{fig:sim_cb}} \\end{figure*}  Ample efforts have been devoted to the accurate determination of the thermal structuring of coronal plasma to derive robust observational constraints on the coronal heating mechanism(s). The plasma temperature distribution of the quiet corona and of active regions has been investigated through imaging data and spectroscopic observations \\citep[e.g.,][]{Brosius96,Landi98,Aschwanden00,Testa02, DZM03,Reale07,Landi09,Shestov10,Sylwester10}. Several recent studies have focused on EUV spectra obtained with the \\hinode\\ Extreme Ultraviolet Imaging Spectrometer (\\eis; \\citealt{Culhane07}) which provides good temperature diagnostic capability, together with higher spatial resolution and temporal cadence than previously available \\citep[e.g.,][]{Watanabe07,Warren08loops,Patsourakos09, Brooks09,Warren09,Testa11,Tripathi11}. Imaging observations obtained by the {\\em SDO} Atmospheric Imaging Assembly (AIA; \\citealt{Lemen12}) in narrow EUV bands have also been used for studying the plasma temperature distribution, especially in active region loops \\citep[e.g.,][]{Schmelz10,Aschwanden11, Aschwanden11b,Schmelz11}, also in conjunction with EIS data \\citep[e.g.,][]{Brooks11,Warren11,Testa12b}.  The optically thin nature of the coronal EUV and X-ray emission implies that the emission observed in the resolution element (instrument pixel) is generally produced by several independent structures along the LOS. Most plasma diagnostics rely on some homogeneity assumptions (e.g., constant plasma density along the LOS) whereas these overlapping structures can in principle have very different plasma parameters. This poses significant challenges for interpreting the meaning of the derived plasma parameters (which are by necessity weighted averages of the distributions along the LOS) in terms of the actual physical conditions of the plasma. Further assumptions typically made for specific diagnostics (e.g., on the functional form and smoothness of the EMD) also have a potentially significant impact on the accuracy of the diagnostics. Several efforts have been carried out to test the accuracy of the plasma temperature diagnostics. For instance, the notorious challenges in determining the emission measure distribution and its confidence limits have been addressed in several studies \\citep[e.g.,][]{Craig76, Judge97,McIntosh00,Judge10,LandiKlimchuk10,Landi12,Hannah12}. The main  limitation of these previous efforts lies in the necessarily simplified test cases adopted.  In this paper we use advanced 3D radiative MHD simulations of the solar atmosphere with the Bifrost code \\citep{Gudiksen11} to carry out detailed tests of coronal plasma temperature diagnostics. We focus on the plasma temperature diagnostics using current spectral ({\\em Hinode/EIS}) and imaging ({\\em SDO/AIA}) data. We synthesize EIS and AIA data from 3D radiative MHD (rMHD) simulations and analyze them like real data, and use the comparison of the thermal distributions inferred from the synthetic data with the ``true'' distributions, which in the case of the simulations are known, to carefully assess the accuracy of the diagnostics. These 3D simulations provide us with the opportunity to improve upon previous work by exploring more realistic configurations, with significant superposition of different structures along the LOS, allowing a statistical approach to determine the accuracy and limitations of the plasma diagnostics, for a variety of spatial and thermal structuring of the plasma. The three dimensional nature of the simulations also allows us to explore a variety of realistic viewing angles, reproducing typical distributions of structuring ranging from on disk to limb observations.  In Section~\\ref{s:method} we describe the analysis methods, we include a short description of the Bifrost code, and discuss the characteristics of the 3D rMHD models and of the corresponding synthetic observables used in this work. The analysis of the synthetic data and the results of the determination of the plasma temperature distribution are presented and discussed in Section~\\ref{s:results}. We summarize our findings and draw our conclusions in Section~\\ref{s:conclusions}. ",
        "conclusions": ""
    },
    "1208/1208.3373_arXiv.txt": {
        "abstract": "We have recently proposed a novel self tuning mechanism to alleviate the famous cosmological constant problem, based on the general scalar tensor theory proposed by Horndeski. The self-tuning model ends up consisting of four geometric terms in the action, with each term containing a free potential function of the scalar field; the four together being labeled as the {\\it  Fab-Four}. In this paper we begin the important task of  deriving the cosmology associated with the {\\it  Fab-Four} Lagrangian. Performing a phase plane analysis of the system we are able to obtain a number of fixed points for the system, with some remarkable new solutions emerging from the trade-off between the various potentials. As well as obtaining inflationary solutions we also find conventional radiation/matter-like solutions, but in regimes where the energy density is dominated by a cosmological constant, and where we do not have any explicit forms of radiation or matter. Stability conditions for matter solutions are obtained and we show how it is possible for there to exist an extended period of `matter domination' opening up the possibility that we can generate cosmological structures, and recover a consistent cosmology even in the presence of a large cosmological constant. ",
        "introduction": "Over the last decade or so, as we have struggled to explain the nature of dark energy that is believed to be responsible for the observed acceleration of the Universe, interest has turned to the possibility that rather than being caused by an unknown form of energy density, the acceleration could be a result of a modification of Einstein's theory of General Relativity. It has resulted in an explosion of papers in the field, see \\cite{review} for a detailed review of the various approaches that have been adopted; one particularly interesting direction involves scalar-tensor combinations. It seems sensible to require that any theory maintains second order field equations, and the most general scalar-tensor theory satisfying that criteria was written down back in 1974 by Horndeski  \\cite{horndeski:1974} (it has recently been rediscovered in \\cite{general}). Such theories of modified gravity cover a wide range of models, ranging from Brans-Dicke gravity \\cite{bdgravity} to the recent models \\cite{covgal,galmodels} inspired by galileon theory \\cite{galileon}; the latter being examples of higher order scalar tensor Lagrangians with second order field equations. Of course all of these models can be considered as special cases of Horndeski's original action. Once the action was rediscovered, it did not take long before a perturbative analysis of the background evolution equations was carried out \\cite{Kobayashi:2011nu,DeFelice:2011hq}, which allows for a stability analysis to be performed on the various background solutions.  In \\cite{Charmousis:2011bf} we obtained a new class of models arising out of Horndeski's theory on FLRW backgrounds. The new models gave a viable self-tuning mechanism for solving the (old) cosmological constant problem, at least at the classical level, by completely screening the spacetime curvature from the net cosmological constant.  In order to evade the famous no-go theorem of  Weinberg \\cite{nogo}, the new solutions did not assume Poincar\\'e invariance to hold at the level of the  solution (as Weinberg assumed), rather we allowed it to be broken in the scalar field sector. This is similar to a route adopted in \\cite{bigalileon} where the scalar field is allowed to break Poincar\\'e invariance on the self-tuning vacua, whilst maintaining a flat spacetime geometry. In \\cite{Charmousis:2011bf} we provided a brief sketch of how the system works, showing that by demanding the self-tuning mechanism continues to work through phase transitions, so causing the vacuum energy to jump,  we get powerful restrictions on the allowed form of Horndeski's original theory. Whereas the  original model is complicated, with many arbitrary functions of both the scalar and its derivatives, we showed that by assuming matter is only minimally coupled to the metric (required to satisfy equivalence principle (EP) considerations) then once the model is passed through our self-tuning filter, it reduces in form to just four base Lagrangians each depending on an arbitrary function of the scalar only, coupled to a curvature term. We called these base Lagrangians {the {\\it  Fab-Four}}: ${\\cal L}_j$, ${\\cal L}_p$, ${\\cal L}_g$, ${\\cal L}_r$, where the indices refer to {\\it John, Paul, George} and {\\it Ringo}. This was followed up in \\cite{Charmousis:2011ea} with a detailed derivation of the conditions that lead to the four base Lagrangians just mentioned, in which we showed how they naturally lead to self-tuning solutions. Moreover in \\cite{Charmousis:2011ea} we began to address the important question of the stability of the classical solutions to quantum corrections and demonstrated that at least heuristically the self-tuning solutions can be guaranteed to receive only small quantum corrections thereby not spoiling the self-tuning nature of the solutions.  The purpose of this paper is to begin the discussion of the cosmology associated with the {\\it  Fab-Four} Lagrangian. Without a sensible cosmology the model is nothing other than an interesting aside that may give some feeling as to how the cosmological constant can be addressed, but in itself does not have anything to say about our Universe. This is a non-trivial exercise.  Note that some aspects of \\FF cosmology were touched upon in \\cite{j+g}.   Recall from \\cite{Charmousis:2011bf} that we are dealing with situations where the net cosmological constant may be large compared to any other energy density in the system. In the conventional cosmological scenario this would inevitably lead to a period of rapid acceleration, with no prospect of a radiation or matter dominated period, hence no chance for nucleosynthesis to take place or structures to form in our Universe. We will require the four potential functions in the {\\it  Fab-Four} action to act together and conspire to alleviate the influence of a net large cosmological constant before the final self tuning solution is reached. After all, we do not live in that solution just yet. To attack the problem, we will rewrite the equations of motion for the scalar field and the Friedmann equation in terms of a dynamical system allowing us to look for late time attractor solutions and to determine the stability of those solutions.  The initial conclusions are positive, we are able to find combinations of the four potentials that do indeed lead to inflation/radiation/matter dominated like periods, with the latter two entering the self-tuning regime at very late times. For matter domination like behaviour, we find that there are  solutions  that are perturbatively stable.  These solutions are remarkable. Because we are interested in the case where the cosmological constant dominates the source, we have focussed on the case where no additional sources are present. Thus, we are able to find perturbatively stable matter-like solutions that are driven by a cosmological constant.  What is happening is that the scalar field is working to screen the pressure component of the cosmological constant before its energy density.  Eventually it will also screen the energy density, but the potentials allow for an intermediate period in which $\\Lambda$ essentially behaves like cold dark matter. This is the main result of this paper.   The layout is as follows: in section \\ref{horndeski-action} we briefly recap the key Hamiltonian and scalar field equations of motion for the {\\it  Fab-Four} system arising from the original action of Horndeski  \\cite{horndeski:1974} that is minimally coupled to matter. We do not rederive them, rather direct the reader to \\cite{Charmousis:2011ea} for a rigorous derivation of the Lagrangian and evolution equations. We begin exploring the cosmology of this self-tuning scenario  in section \\ref{sec-individual}, focussing on how each member of the \\FF behaves in isolation. To see how the various members behave in combination we rewrite the field equations as first order equations using a dynamical systems approach in section \\ref{self-tune-cosmology}.  We switch off curvature in order to focus on the cosmological epoch prior to self-tuning, and find scaling solutions corresponding to different types of cosmology such as radiation domination, matter domination and inflation. Strictly speaking, some of these matter-like solutions are only fixed points for vanishing cosmological constant.  Even so, as we show in section \\ref{sec-summary}, both analytically and through numerical simulations, they still provide an excellent approximation to the true cosmology even when there is a large non-vanishing cosmological constant.  The reader not overly concerned with the details of how we arrived at interesting classes of \\FF potentials should probably skip section \\ref{self-tune-cosmology} and proceed directly to section \\ref{sec-summary}.  Here we summarize the main findings of section \\ref{self-tune-cosmology}, as well as providing numerical simulations of solutions when spatial curvature is turned on.  For the matter-like solutions, we also consider cosmological perturbations to weed out any problems with ghost and/or gradient instabilities. Whilst some solutions are unstable, others are  perfectly well behaved. We conclude in section \\ref{conc}. ",
        "conclusions": "\\label{conc} Obtaining a sensible cosmology in the presence of a large and changing contribution to the vacuum energy $\\rho_\\Lambda$ is one of major challenges facing the self tuning scenario that we have developed in \\cite{Charmousis:2011bf,Charmousis:2011ea}. A standard cosmological evolution arising out of General Relativity with large $\\rho_\\Lambda$ would be totally unacceptable apart from perhaps in the early Universe where it would drive a period of accelerated expansion. There would be no way of exiting this period of inflation and obtaining a period of radiation and matter domination required for nucleosynthesis and structure formation. The goal of this paper is to demonstrate that for the \\FF,  it is indeed possible to obtain  a sensible cosmological history even in the presence of a large $\\rho_\\Lambda$ contribution at all times. In other words, the self tuning of the cosmological constant can be accommodated in a sensible cosmological timeline.  To show this we have developed a dynamical systems approach in which fixed point solutions corresponding to inflationary, radiation and matter dominated solutions are made manifest. Two key approaches are developed. In the first, we explicitly include the $\\rho_\\Lambda$ contribution and by demanding that all contributions in the Hamiltonian constraint remain independent of the Hubble parameter, we show that there exist a class of scaling solutions corresponding to the cosmologies we are looking for (recall there is no requirement here to actually include matter or radiation sources, the scalar field is doing the work for us). However, it turns out that this particular matter dominated solution (called ``Matter III\"), whilst perfectly acceptable at the background level, actually contains a gradient instability when perturbed, an instability that would grow on too fast a timescale to be compatible with observations. This has led us to consider a second complementary approach. Rather than obtain background solutions in the presence of $\\rho_\\Lambda$ we set it to zero and look for consistent solutions that can also mimic  matter domination ($H^2 \\propto a^{-3}$). Our requirement that the Hamiltonian constraint be independent of the Hubble parameter is now lifted and this allows for more freedom introducing an extra parameter (we call $n$ or equivalently, $\\hat n $) in the background solutions.  We find  new classes of matter-like solutions (called ``Matter I\" , `Matter II\" and ``Matter IV\") .  It would seem to go against the self-tuning spirit of the \\FF that the ``Matter I\" and ``Matter II\" solutions only correspond to fixed points for vanishing cosmological constant. However, we have shown that they can still represent an excellent approximation even when a large $\\rho_\\Lambda$ is turned on, provided $\\hat n>0$.  This is because the solution gets corrected by $\\rho_\\Lambda/a^{\\hat n}$, a correction that decreases with time as the scale factor grows.   Using analytic methods,  we also showed  that for vanishing spatial curvature,  these solutions are cosmological attractors for $\\hat n>0$. Once spatial curvature is reintroduced alongside the cosmological constant, we are forced to use numerical simulations which reproduce the expected behaviour: a long period of matter-domination, before asymptoting to the self-tuning Milne Universe.  In contrast to the  Matter I, II, and III potentials, the Matter IV potentials have  corresponding solutions that are stable against cosmological perturbations.  This opens  up the possibility of a sensible matter dominated period of evolution, hence of structure formation in the {\\it  Fab-Four} scenario.  Furthermore, these solutions are behaving in such a way that the scalar screens the pressure component of the cosmological constant before the energy density. At least for homogeneous and isotropic backgrounds, this suggests that the cosmological constant is being forced to behave like cold dark matter.  It is certainly tempting to ask whether such behaviour extends to inhomogeneous solutions, and recent results suggest that it may well be possible to have a \\FF scenario satisfying current solar system constraints \\cite{Rinaldi:2012vy}.  There is much that remains to be done. We have not yet obtained a full cosmology, but the fact that we have a class of background polynomial potentials that we know can provide the various cosmological epochs we want to reproduce offers us some direction.  Indeed we can speculate as to how we might sew together these interesting potentials to achieve the desired results. The point is that the scalar field is continually evolving, so we could arrange for the potential to correspond to different fluid behaviours for different ranges of $\\phi$. For example, if we want radiation domination for $H^2>H^2_{eq}$ and  ``Matter I\" like behaviour for $H^2<H^2_{eq}$, we might propose a Lagrangian of the form \\be {\\cal L}= \\frac{m^4}{H_{eq}^2} \\phi^{\\hat n} R+ {\\cal L}_{\\textrm{``Matter I\"}}\\Big |_{c_1 \\sim \\order(1), c_2 \\sim \\order(m^{-2})}, \\ee Plugging in the ``Matter I' solution (\\ref{sol1}), we see that \\be \\frac{m^4}{H_{eq}^2} \\phi^{\\hat n} R \\sim m^4 a^{\\hat n} \\left(\\frac{a_{eq}}{a} \\right)^3, \\qquad {\\cal L}_{\\textrm{``Matter I\"}}\\Big |_{c_1 \\sim \\order(1), c_2 \\sim \\order(m^{-2})} \\sim  m^4 a^{\\hat n}, \\ee where $a_{eq}=(m/H_{eq})^{2/3}$ is the value of the scale factor when $H=H_{eq}$. We see that the ``Matter I\" terms dominate for $a>a_{eq}$, as desired for mater domination.  For $a<a_{eq}$ the $\\phi^{\\hat n}R$ term becomes important, and might be expected to dominate the dynamics, yielding an earlier  period of radiation domination. The same may be done to evade the graceful exit problem due to the inflationary solutions being attractors."
    },
    "1208/1208.4848_arXiv.txt": {
        "abstract": "{The APEX Telescope Large Area Survey: The Galaxy (ATLASGAL) is an unbiased continuum survey of the inner Galactic disk at 870 $\\mu$m. It covers $\\pm 60^{\\circ}$ in Galactic longitude and aims to find all massive clumps at various stages of high-mass star formation in the inner Galaxy, particularly the earliest evolutionary phases.} {We aim to determine properties such as the gas kinetic temperature and dynamics of new massive cold clumps found by ATLASGAL. Most importantly, we derived their kinematical distances from the measured line velocities.} {We observed the ammonia ($J$,$K$) = (1,1) to (3,3) inversion transitions toward 862 clumps of a flux-limited sample of submm clumps detected by ATLASGAL and extracted $^{13}$CO (1-0) spectra from the Galactic Ring Survey (GRS). We determined distances for a subsample located at the tangential points (71 sources) and for 277 clumps whose near/far distance ambiguity is resolved.} {Most ATLASGAL clumps are cold with rotational temperatures from 10-30 K with a median of 17 K. They have a wide range of NH$_3$ linewidths (1-7 km~s$^{-1}$) with 1.9 km~s$^{-1}$ as median, which by far exceeds the thermal linewidth, as well as a broad distribution of high column densities from $10^{14}$ to $10^{16}$ cm$^{-2}$ (median of $2 \\times 10^{15}$ cm$^{-2}$) with an NH$_3$ abundance in the range of 5 to $30 \\times 10^{-8}$. ATLASGAL sources are massive, $\\gtrsim$ 100 M$_{\\odot}$, and a fraction of clumps with a broad linewidth is in virial equilibrium. We found an enhancement of clumps at Galactocentric radii of 4.5 and 6 kpc. The comparison of the NH$_3$ lines as high-density probes with the GRS $^{13}$CO emission as low-density envelope tracer yields broader linewidths for $^{13}$CO than for NH$_3$. The small differences in derived clump velocities between NH$_3$ (representing dense core material) and $^{13}$CO (representing more diffuse molecular cloud gas) suggests that the cores are essentially at rest relative to the surrounding giant molecular cloud.} {The high detection rate (87\\%) confirms ammonia as an excellent probe of the molecular content of the massive, cold clumps revealed by ATLASGAL. A clear trend of increasing rotational temperatures and linewidths with evolutionary stage is seen for source samples ranging from 24 $\\mu$m dark clumps to clumps with embedded HII regions. The survey provides the largest ammonia sample of high-mass star forming clumps and thus presents an important repository for the characterization of statistical properties of the clumps and the selection of subsamples for detailed, high-resolution follow-up studies.} ",
        "introduction": "High-mass stars are essential for the evolution of galaxies because they energize the interstellar medium and release heavy elements, which are determining Galactic cooling mechanisms. Although the understanding of high-mass star formation is consequently of great importance, only little is known about the early stages of the formation of massive stars in contrast to the well-established evolutionary sequence of isolated low-mass stars \\citep{2000prpl.conf...59A}. This is partly because of the difficult observational conditions that one encounters. High-mass stars form in clusters and are rarer according to the initial mass function \\citep{2011arXiv1112.3340K}, therefore they are usually found at larger distances than molecular clouds, in which low-mass stars form. Moreover, high-mass stars evolve on shorter time scales and within dynamic and complex environments and are still deeply embedded in their earliest phases.\\\\ It is well known that high-mass star formation starts in giant molecular clouds \\citep{2001ApJ...547..792D}. In analogy to low-mass star formation, one expects a cold and massive starless clump of $\\gtrsim$ 500 M$_{\\odot}$ consisting of molecular gas for the initial conditions of massive star/cluster formation \\citep{2007prpl.conf..165B}. Young massive (proto) stars evolving in dense cores within the clumps emit ultraviolet radiation that leads to the formation of ultracompact HII regions (UCHIIRs). Many of them were found in radio continuum surveys with the VLA \\citep[e.g.][]{1989ApJS...69..831W}, which revealed their distribution and number. Those regions were very helpful in tracing a more evolved stage of massive star formation in the Galaxy. Often UCHIIRs are associated with so-called hot molecular cores \\citep{1992A&A...256..618C}, which are possible precursors of UCHIIRs, also called high-mass protostellar objects (HMPOs) or massive young stellar objects (MYSOs), which have also been discovered recently \\citep{2002ApJ...566..945B}. Still earlier phases are deeply embedded in a massive envelope and are so cold that mostly they cannot be probed by means of mid-infrared emission \\citep[e.g.][]{2005ApJ...634L..57S, 2010ApJ...723..555P,2007A&A...476.1243M}. Under certain circumstances, these clumps can be seen in absorption against a bright background provided mostly by emission from polycyclic aromatic hydrocarbon molecules (PAHs) as infrared dark clouds \\citep[IRDCs,][]{1996A&A...315L.165P, 1998ApJ...494L.199E}. \\cite{2006ApJ...653.1325S} and \\cite{2010ApJ...723..555P} have presented large catalogues of these IRDCs. But to date, only a small fraction of them has been investigated in detail \\citep{1998ApJ...508..721C,2000ApJ...543L.157C,2006A&A...450..569P,2006ApJ...641..389R}. Previous studies revealed that IRDCs are cold ($<$ 25 K), dense ($>$ $10^5$ cm$^{-3}$) and have high column densities \\citep[$\\sim$ 10$^{23}$-10$^{25}$ cm$^{-2}$,][]{2006ApJ...641..389R,2006ApJ...639..227S}. Their sizes ($\\sim$ 2 pc) and masses ($\\sim$ 10$^2$-10$^4$ M$_{\\odot}$) are similar to those of cluster-forming molecular clumps. This, together with the cold temperatures and fragmented substructure of IRDCs, suggests that they are protoclusters \\citep{2006ApJ...641..389R,2005ApJ...630L.181R}. Kinematic distances to IRDC samples in the first and fourth Galactic quadrants have been determined using $^{13}$CO (1-0) and CS (2-1) emission \\citep{2006ApJS..163..145J,2008ApJ...680..349J}. However, surveys targeted at the early phases of high-mass star formation \\citep[e.g.][]{2002ApJ...566..945B,1989ApJ...340..265W} usually only trace one of those stages, e.g. IRDCs, HMPOs or UCHIIRs. In contrast, a first attempt to obtain an unbiased sample using a complete dust continuum imaging at the scale of the molecular complex Cygnus X \\citep{2007A&A...476.1243M} was able to derive some more systematic results on the existence of a cold phase (IR-quiet massive dense cores and Class 0-like massive protostars) for high-mass star formation \\citep{2010A&A...524A..18B,2011IAUS..270...53C}. Another survey of the molecular complex NGC 6334/NGC 6357 was conducted by \\cite{2010A&A...515A..55R}, who identified IR-quiet massive dense cores and estimated high-mass protostellar lifetimes. The IR-quiet massive dense cores are suggested to host massive class 0 protostars as seen by the Herschel Space Telescope \\citep{2011A&A...535A..76N}.\\\\ Only a few hundred high-mass proto- or young stellar objects have been studied up to now. Hence, these and especially objects in still earlier phases of massive star formation need to be investigated in more detail. To achieve significant progress in that, a large-area survey of the Galactic plane was conducted, which now provides a global view of star formation at submillimeter wavelengths. ATLASGAL (\\textit{The APEX Telescope Large Area Survey of the Galaxy at 870 $\\mu$m}) is the first unbiased continuum survey of the whole inner Galactic disk at 870 $\\mu$m \\citep{2009arXiv0903.1369S}. It aims to find all massive clumps that form high-mass stars in the inner Galaxy by using the Large APEX Bolometer Camera LABOCA, which is an array with 295 bolometer pixels operated at the APEX telescope \\citep{2007Msngr.129....2S, 2008SPIE.7020E...2S} with a field of view of 11$\\arcmin$ and a beamwidth of 19.2$\\arcsec$ FWHM at the wavelength of 870 $\\mu$m.\\\\ ATLASGAL, which reaches a Galactic longitude of $\\pm 60^{\\circ}$ and latitude of $\\pm 1.5^{\\circ}$, is able to detect objects associated with massive star formation at various stages and to compare them \\citep{2009arXiv0903.1369S}. Other Galactic plane surveys, conducted over a similar range, are the MSX survey covering 8 to 21 $\\mu$m \\citep{2001AJ....121.2819P}, 2MASS around 2 $\\mu$m \\citep{2006AJ....131.1163S}, GLIMPSE from 3 to 8 $\\mu$m \\citep{2003PASP..115..953B}, MIPSGAL at 24 and 70 $\\mu$m \\citep{2009PASP..121...76C} and HiGAL from 70 to 500 $\\mu$m \\citep{2010PASP..122..314M}.  While submillimeter dust continuum surveys are essential for identifying high-mass clumps, they lack information on important parameters, especially the distances to the newly found sources. These are needed to determine other properties such as their masses and luminosities. But all of this can be addressed by molecular line observations of the high-mass star forming regions. Because their density is at about the critical NH$_3$ density, this molecule is appropriate for determining properties of the clumps without contamination from large-scale cloud structures with lower density. Since the molecular gas of the cores is very dense \\citep[$\\sim$ 10$^5$ cm$^{-3}$][]{2002ApJ...566..945B} and cold, many molecules such as CS and CO are partly frozen onto dust grains. In contrast to those, the fractional gas abundance of ammonia remains constant in prestellar cores \\citep{2002ApJ...569..815T}. Hence, we carried out follow-up observations of northern ATLASGAL sources in the lowest NH$_3$ inversion transitions.\\\\ NH$_3$ is known as a reliable temperature probe of interstellar clouds \\citep{2004A&A...416..191T,1983A&A...122..164W}. The rotational energy levels are given by the total angular momentum, $J$, and its projection along the molecular axis, $K$. Radiative transitions between different $K$-ladders are forbidden and the lowest metastable energy levels, for which $J = K$, are thus collisionally excited. The intensity ratio of their inversion transitions therefore provides the rotational temperature of the gas  \\citep{1983ARA&A..21..239H}, which can be used to estimate the kinetic temperature of the cores. This is needed for proper mass estimates from the submm data. Independently, the linewidth can be used to derive virial masses. Moreover, the inversion transitions are split into distinct hyperfine components and their ratio provides a measure of the optical depth, knowledge of which leads to reliable column density and rotational temperature determinations.\\\\ As this article was ready for submission, we became aware of the recent study by \\cite{2011ApJ...741..110D}, who observed NH$_3$ lines of sources in the first Galactic quadrant as well, identified by the Bolocam Galactic Plane Survey (BGPS), which measures the 1.1 mm continuum emission. In contrast to our sample, the observations of \\cite{2011ApJ...741..110D} do not cover the whole northern Galactic longitude range up to 60$^{\\circ}$, but are conducted only within four different ranges in longitude, and within a smaller Galactic latitude range of $\\pm 0.5^{\\circ}$ compared to our NH$_3$ survey. Our analysis is thus complementary to the results of \\cite{2011ApJ...741..110D}.\\\\ In Section \\ref{s:obs}, we describe which ATLASGAL sources were selected for our NH$_3$ follow-up observations and provide details of the measurements. In Section \\ref{data reduction}, we present how we reduced the data and derived different clump properties from fits to the spectra of NH$_3$ inversion transitions. In Section \\ref{results}, different correlations of the NH$_3$ line parameters such as the velocity distribution, linewidth, rotational temperature and column density are shown. Then, we compare these with submillimeter dust continuum properties of the clumps such as the H$_2$ column density in Section \\ref{dust continuum} and estimate gas and virial masses. We analyse additional line parameters derived from $^{13}$CO (1-0) emission \\citep{2006ApJS..163..145J} in Section \\ref{13CO} to investigate clump-to-cloud motions etc. In Section \\ref{discussion}, the determined gas properties are compared with those of other high-mass dust-selected star forming regions. A summary and conclusions of this NH$_3$ investigation are given in Section \\ref{conclusion}.\\\\ This paper analyses statistics of the NH$_3$ data obtained for the northern ATLASGAL sources. We provide their near and far distance, but do not distinguish between both, although some of the molecular clouds have known distances from a previous study \\citep{2009ApJ...699.1153R}. In a second paper (Wienen et al. in prep.) ammonia lines of ATLASGAL sources in the southern hemisphere will be studied and distances to southern molecular clouds will be derived. ",
        "conclusions": "\\label{conclusion} The (1,1) to (3,3) ammonia inversion transitions of 898 ATLASGAL sources, located between 5$^{\\circ}$ and $60^{\\circ}$ in Galactic longitude and $\\pm 1.5^{\\circ}$ in Galactic latitude, were observed with the Effelsberg telescope. They build the largest available ammonia sample of high-mass star forming clumps. Our results are summarised in this section.\\\\ We analysed the correlation of different ammonia line parameters and compared them to the submm dust continuum and $^{13}$CO emission.\\\\ The NH$_3$ (1,1) line's LSR velocity of most clumps lies between -10 km~s$^{-1}$ and 120 km~s$^{-1}$, most clumps show a strong correlation with CO emission \\citep{2001ApJ...547..792D}, probing the larger giant molecular clouds. The velocities yield near/far kinematical distances, that were calculated for 750 unknown sources using the revised rotation parameters of the Milky Way presented by \\cite{2009ApJ...700..137R}. Galactocentric distances show an enhancement of ATLASGAL sources at 4.5 and 6 kpc, tracing the Scutum arm and Sagittarius arm, which agrees with previous results for other tracers of massive star formation \\citep{2009ApJ...690..706A}.\\\\ Most sources exhibit NH$_3$ (1,1) linewidths between 0.7 km~s$^{-1}$ and 3.5 km~s$^{-1}$, some range up to 7 km~s$^{-1}$, which is much broader than the thermal NH$_3$ linewidth and common linewidths observed in cores of low-mass \\citep{1999ApJS..125..161J}. The rotational temperature between the (1,1) and (2,2) inversion levels was derived and was used to calculate the kinetic temperature. Most ATLASGAL clumps have rotational temperatures between 10 K and 25 K, some exhibit up to 30 K, while their average kinetic temperatures range from 12 K to 35 K, only few up to 40 K. The increasing rotational temperature and the broader widths of the (3,3) line, which probes higher temperature, hint at star formation activity inside the clumps. Indeed, many clumps show a stronger than expected intensity of the NH$_3$ (3,3) line as estimated for a homogeneous, cold clump with a temperature derived from the (1,1) and (2,2) lines. This results probably from an embedded hot core component with a low beam filling factor and high temperature.   \\\\ We used the kinetic temperature and submm dust continuum flux to determine the H$_2$ density. Its correlation with the NH$_3$ column density leads to ammonia abundances relative to H$_2$ mostly from 5$\\times$10$^{-8}$ to 3$\\times$10$^{-7}$.\\\\ For ATLASGAL clumps located at tangential points and associated with clouds from the GRS survey, preliminary distance estimates were given and gas masses were calculated, that lie in the range from 60 to 10$^4$ M$_{\\odot}$. Comparing them with the virial masses yields a virial parameter of $\\sim$ 1 for sources with broad linewidths, which are supported against gravitational collapse. In contrast, clumps with narrow linewidths have a virial parameter of only $\\sim$ 0.5.\\\\ The comparison of the NH$_3$ lines as high-density probe with the GRS $^{13}$CO emission as low-density envelope tracer yields broader $^{13}$CO than ammonia linewidths. This can result from turbulence within large structures probed by $^{13}$CO. We also analysed the relative motion of the clump and its cloud using NH$_3$, which represents the dense core material and $^{13}$CO to trace the more diffuse molecular cloud gas. Since our comparison yields small differences between NH$_3$ and $^{13}$CO velocities, less than the average $^{13}$CO linewidth, random motions of the ammonia clumps are expected and hint at turbulence. In contrast to low-mass cores, which show mainly subsonic core-to-envelope motions \\citep{2004ApJ...614..194W, 2007ApJ...655..958W, 2007ApJ...668.1042K}, the relative motions of our sample mostly exceed the sound speed.\\\\ Recent NH$_3$ (1,1) to (3,3) line observations of northern BGPS sources were carried out by \\cite{2011ApJ...741..110D} at the Robert F. Byrd Green Bank Telescope (GBT) with a lower FWHM of 31$\\arcsec$ compared to the Effelsberg beamwidth. Our results agree well with those of \\cite{2011ApJ...741..110D}, yielding similar ranges of kinetic temperature, NH$_3$ column density, H$_2$ column density, thus NH$_3$ abundance, and gas masses.\\\\ Association of the ATLASGAL sources with other high-mass star forming samples led to subsamples in various evolutionary stages, whose NH$_3$ line parameters were compared. Clumps with low temperatures, narrow linewidths, and high column densities may be IRDCs, which are assumed to be the cold precursors to hot cores and are consequently in an early evolutionary phase of high-mass star formation \\citep{2006ApJ...641..389R}. Moreover, few ATLASGAL sources show higher temperatures, greater linewidths and lower column densities than the IRDCs. Those properties hint at objects in a more evolved stage of massive star formation, such as YSOs that contain a luminous infrared source or clumps with embedded UCHII regions.\\\\ Because these NH$_3$ observations and analyses can also build a valuable database for the selection of different phases of massive star formation and the determined velocities are important especially for the so far unknown ATLASGAL sources, we extended the NH$_3$ survey to the south using the Parkes telescope. The comparison of first and fourth quadrant properties and the distinction between near and far distances will be published in a forthcoming paper (Wienen et al. in prep.).\\\\"
    },
    "1208/1208.5143_arXiv.txt": {
        "abstract": "Using new and published data, we construct a sample of \\nsample\\ brightest cluster galaxies (BCGs) spanning the redshift interval $0.03 < z < 1.63$. We use this sample, which covers 70\\% of the history of the universe, to measure the growth in the stellar mass of BCGs after correcting for the correlation between the stellar mass of the BCG and the mass of the cluster in which it lives. We find that the stellar mass of BCGs increase by a factor of \\growth\\ between $z=0.9$ and $z=0.2$. Compared to earlier works, our result is closer to the predictions of semi-analytic models. However, BCGs at $z=0.9$, relative to BCGs at $z=0.2$, are still a factor of 1.5 more massive than the predictions of these models. Star formation rates in BCGs at $z\\sim1$ are generally too low to result in significant amounts of mass. Instead, it is likely that most of the mass build up occurs through mainly dry mergers in which perhaps half of the mass is lost to the intra-cluster medium of the cluster. ",
        "introduction": "Brightest Cluster Galaxies (BCGs) are amongst the largest, most luminous and most massive galaxies in the universe at the present epoch. Located in the cores of rich galaxy clusters, BCGs are easy to identify, both observationally and in simulations. They can also be observed at a time when the universe was less than a third of its current age. They therefore provide an attractive target for testing our understanding of the processes that drive galaxy evolution, albeit in the most massive galaxies of the universe.  In the hierarchical scenario for the formation of structure in our universe, galaxies start off as small fluctuations in the density of matter and build up their stellar mass over time by converting material accreted from their surroundings into stars and by merging with other galaxies \\citep[see][for a review]{Baugh2006}. In semi-analytic models that use the hierarchical scenario as their foundation, the stellar mass of a BCG increases significantly with time. For example, between redshift $z=1.0$ (corresponding to a look-back time of 6.7\\,Gyr) to $z=0$, the semi-analytic model described in \\citet{DeLucia2007} predicts that BCGs increase their stellar mass by a factor of four \\citep{DeLucia2007}.  In contrast to this prediction, observations appear to suggest that there is little growth in the stellar mass of BCGs, although apparently conflicting results have been reported. Using a sample of optically selected clusters, \\citet{Aragon-Salamanca1998} found that the stellar mass of BCGs grew by a factor of 4 between $z=1$ and today. \\citet{Burke2000}, on the other hand, using a sample of X-ray selected clusters over a similar redshift range, find substantially less growth.\\footnote{Both \\citet{Aragon-Salamanca1998} and \\citet{Burke2000} use an Einstein de-Sitter universe, i.e.~$\\Omega_{\\mathrm M},\\Omega_{\\Lambda}=1,0$, with $H_0=50$ km\\,s$^{-1}$\\,Mpc$^{-1}$ for the cosmology. While their results are not directly comparable to the results in later papers, one can compare the results of the two papers.}  \\citet{Burke2000} conclude that sample selection can explain part of the difference between their results and those in \\citet{Aragon-Salamanca1998}, a conclusion that was supported by \\citet{Nelson2002}. In an independent study, using an optically selected sample of 21 high-redshift clusters, \\citet{Whiley2008} find little change in the stellar mass of BCGs since $z\\sim 1$.  At higher redshifts, the discrepancy between the models and the observations is larger. \\citet{Collins2009} and \\citet{Stott2010}, using a sample of 20 mostly X-ray selected clusters and a sample of nearby clusters from \\citet{Stott2008}, find that there is little growth between $z \\sim 1.4$ and now. At $z \\sim 1.4$, the semi-analyic model of \\citet{DeLucia2007} predicts that BCGs should be a factor of six less massive. Therefore, there appears to be clear disagreement between the models and the observations.  In this paper, we expand upon the work that has been done so far in three ways. First, we increase the number of BCGs beyond $z=0.8$ for which accurate near-IR photometry is available. Second, we extend the redshift baseline by including the BCGs in two recently discovered clusters at $z\\sim 1.6$. Third, we use our large sample to account for the correlation between the stellar mass of the BCG and the mass of the cluster in which it lives.  We start the paper with a description of our new sample of BCGs in Section 2, followed, in Section 3, with a description of the near-IR imaging data that we use in later sections. In Sections 4 and 5, we derive the magnitudes and colours of the BCGs in our sample and compare them to predictions made by simple and composite stellar population models.  Following \\citet{Stott2010}, we use this comparison to estimate stellar masses. In Section 6, we discuss our results, comparing them to the predictions made by semi-analytic models and examining how robust they are to our methods.  In the final section, we summarise our main results.  Throughout the paper, all magnitudes and colours are measured in the observer frame and are placed on the 2MASS photometric system. Vega magnitudes are used throughout the paper. We also assume a flat $\\Lambda$CDM cosmology with $\\Omega_{\\Lambda}=0.73$ and $H_0=70$ km\\,s$^{-1}$\\,Mpc$^{-1}$. % ",
        "conclusions": "\\centering \\begin{tabular}{l r r r l r r} \\hline Cluster & Redshift & R.A.$^5$  & Decl. $^5$ & Instrument/Telescope &\\multicolumn{2}{c}{Exposure times: J and Ks}\\\\ &          & J2000     & J2000      &                     &  [s] & [s] \\\\ \\hline SpARCS J003442-430752\\,$^1$     & 0.867 & 00:34:42.03 & -43:07:53.4  & ISPI/Blanco            & 17280 & 8800 \\\\ SpARCS J003645-441050\\,$^1$     & 0.869 & 00:36:44.99 & -44:10:49.8  & ISPI/Blanco            & 17280 & 7080 \\\\ SpARCS J161314+564930\\,$^{1,2}$  & 0.871 & 16:13:14.63 &  56:49:30.0  & WIRCAM/CFHT            & 6240  & 6300 \\\\ SpARCS J104737+574137\\,$^1$     & 0.956 & 10:47:33.43 &  57:41:13.4  & WIRCAM/CFHT            & 7560  & 2400 \\\\ SpARCS J021524-034331\\,$^1$     & 1.004 & 02:15:23.99 & -03:43:32.2  & ISPI/Blanco            & 26640 & 11800 \\\\ SpARCS J105111+581803\\,$^1$     & 1.035 & 10:51:11.22 &  58:18:03.3  & WIRCAM/CFHT            &  6840 & 2700 \\\\ SpARCS J161641+554513\\,$^{1,2}$  & 1.156 & 16:16:41.32 &  55:45:12.4  & WIRCAM/CFHT            & 18960 & 7000 \\\\ SpARCS J163435+402151\\,$^{1,3}$  & 1.177 & 16:34:38.21 & 40:20:58.4  & WIRCAM/CFHT            & 11640 & 6850 \\\\ SpARCS J163852+403843\\,$^{1,3}$  & 1.196 & 16:38:51.64 &  40:38:42.8  & WIRCAM/CFHT            & 11640 & 6000 \\\\ SpARCS J003550-431224\\,$^{1,4}$  & 1.335 & 00:35:49.68 & -43:12:23.8 & HAWK-I/Yepun  & 11040 & 12000 \\\\ \\hline Cluster & Redshift & R.A.$^5$  & Decl. $^5$ & Instrument/Telescope &\\multicolumn{2}{c}{Exposure times: Y and Ks}\\\\ &          & J2000     & J2000      &                     &  [s] & [s] \\\\ \\hline SpARCS J033056-284300          & 1.626 & 03:30:55.87 &-28:42:59.7  & HAWK-I/Yepun   & 8880 & 3040 \\\\ SpARCS J022426-032331    & 1.633 & 02:24:26.32 & -03:23:30.7 & HAWK-I/Yepun  & 8640 & 5040 \\\\   \\hline \\multicolumn{7}{l}{Note 1: \\citet{Muzzin2012}} \\\\ \\multicolumn{7}{l}{Note 2: \\citet{Demarco2010a}}\\\\ \\multicolumn{7}{l}{Note 3: \\citet{Muzzin2009a}} \\\\ \\multicolumn{7}{l}{Note 4: \\citet{Wilson2009a}} \\\\ \\multicolumn{7}{l}{Note 5: Coordinates of the BCG} \\\\  %  \\end{tabular} \\end{table*}   \\begin{table*} \\caption{Instrument summary}\\label{tab:Instruments} \\label{table:summary} \\centering \\begin{tabular}{l l r r l} \\hline Instrument & Telescope & Pixel Scale & FoV  & Detector\\\\ & & [\\arcsec] & [\\arcmin] & \\\\ \\hline WIRCAM${^1}$ & CFHT & 0.304 & 20.5 & 2x2 Hawaii-2RG mosaic\\\\ ISPI${^2}$ & Blanco & 0.307& 10.3 & Hawaii-2 \\\\ HAWK-I${^3}$ & Yepun (VLT-UT4) & 0.1065& 7.5 & 2x2 Hawaii-2RG mosaic\\\\ \\hline \\multicolumn{5}{l}{Note 1: \\citep{Puget2004a}}\\\\ \\multicolumn{5}{l}{Note 2: \\citep{vanderBliek2004}} \\\\ \\multicolumn{5}{l}{Note 3: \\citep{Pirard2004a,Casali2006a}}\\\\ \\end{tabular} \\end{table*}   \\subsection{Data Reduction}  The processing of the raw data was done in a standard manner and largely follows the steps outlined in \\citet{Lidman2008a}.  Data from each of the cameras were pre-processed (dark subtraction, flat-fielding, and sky subtraction) using a combination of observatory-developed instrument pipelines (for example, the CFHT data were processed with version 1.0 of the `I`iwi pipeline\\footnote{http://www.cfht.hawaii.edu/Instruments/Imaging/WIRCam/}) and our own scripts using IRAF\\footnote{IRAF is distributed by the National Optical Astronomy Observatories which are operated by the Association of Universities for Research in Astronomy, Inc., under the cooperative agreement with the National Science Foundation}.  SCAMP (version 1.6.2) and SWarp (version 2.17.6)\\footnote{http://www.astromatic.net/} were used to map the sky-subtracted images onto a common astrometric reference frame. After accounting for gain variations between chips (only relevant for the data that were taken with HAWK-I and WIRCAM) and creating individual bad pixel maps to account for bad pixels and remnants from bright stars observed in previous frames, the images were then combined with the {\\tt imcombine} task within IRAF.  Each image was weighted with the inverse square of the FWHM of the PSF.  With the exception of the data taken in the Y band, zero points were set using stars from the 2MASS point source catalogue \\citep{Skrutskie2006a}. Typically, between 10 to 40 unsaturated 2MASS stars with 2MASS quality flags of 'A' or 'B' were selected to measure zero points and their uncertainties.  2MASS stars were weighted by the reported uncertainties in the 2MASS point source catalogue. The uncertainties in the zero points are generally less than 2\\%, and more typically 1\\%, for both J and Ks. For Y, the zero point was set using standard stars that were observed during the same night as the clusters. The uncertainty is estimated from the night-to-night variation in the zero points and is around 2\\%.  \\subsection{Data quality}  Overall, the depth and quality of the imaging data varies substantially from one image to another.  The image quality, as measured from bright stars, varies from 0\\farcs3 in the data taken with HAWK-I to 1\\farcs5 in the data taken with ISPI.  The image depth, which we define as the 5 sigma point source detection limit, varies from 19.5 for the Ks band image of SpARCS J003645-441050 to 25.1 for the Y band image of SpARCS J022426-032331. In all cases, the BCG is at least 2 mag brighter than the detection limit. The median signal-to-noise ratio is around 50. Table~\\ref{table:datasummary} summarises the image quality and image depth.   \\begin{table*} \\caption{Image quality and image depth} \\label{table:datasummary} \\centering \\begin{tabular}{l c c || c c} \\hline Cluster & Image quality & Image depth$^1$ & Image quality & Image depth$^1$ \\\\ & [\\arcsec]     & [mag]           & [\\arcsec]     & [mag]           \\\\ \\hline & \\multicolumn{2}{c}{J} & \\multicolumn{2}{c}{Ks}\\\\  \\hline SpARCS J003442-430752  &  1.25 & 21.8  &  0.98  & 19.9 \\\\ % SpARCS J003645-441050  &  1.13 & 21.7  &  1.47  & 19.5 \\\\ % SpARCS J161314+564930  & 0.77 & 22.2  &  0.72  & 21.1 \\\\ SpARCS J104737+574137  & 0.69 & 22.2  &  0.60  & 21.2 \\\\ SpARCS J021524-034331  & 1.07 & 21.8  &  0.89  & 20.3 \\\\ % SpARCS J105111+581803  & 0.66 & 22.5  & 0.74  & 20.6 \\\\ SpARCS J161641+554513  & 0.70 & 22.8  & 0.75  & 21.2 \\\\ SpARCS J163435+402151  & 0.65 & 22.9  & 0.67  & 21.2 \\\\ SpARCS J163852+403843  & 0.61 & 23.1  & 0.58  & 21.5 \\\\ SpARCS J003550-431224   & 0.35 & 24.6  & 0.31  & 23.1 \\\\ \\hline & \\multicolumn{2}{c}{Y} & \\multicolumn{2}{c}{Ks}\\\\ \\hline SpARCS J033056-284300  &   0.45 & 24.0 &  0.29 & 21.9\\\\ SpARCS J022426-032331  &   0.34 & 25.1 &  0.51 & 21.5\\\\   \\hline \\end{tabular} \\medskip \\begin{minipage}{156mm} $^1$The image depth is the 5 sigma point-source detection limit measured over an aperture that has a diameter that is twice the image quality. \\end{minipage} \\end{table*}"
    },
    "1208/1208.0599_arXiv.txt": {
        "abstract": "We present a \\textit{Herschel} far-infrared study towards the rich massive star-forming complex G305, utilising PACS 70, 160\\,\\micron\\,\\,and SPIRE 250, 350, and 500\\,\\micron\\,\\,observations from the Hi-GAL survey of the Galactic plane. The focus of this study is to identify the embedded massive star-forming population within G305, by combining far-infrared data with radio continuum, H$_{2}$O maser, methanol maser, MIPS, and Red MSX Source survey data available from previous studies. By applying a frequentist technique we are able to identify a sample of the most likely associations within our multi-wavelength dataset, that can then be identified from the derived properties obtained from fitted spectral energy distributions (SEDs). By SED modelling using both a simple modified blackbody and fitting to a comprehensive grid of model SEDs, some 16 candidate associations are identified as embedded massive star-forming regions. We derive a two-selection colour criterion from this sample of log\\,(F$_{70}$/F$_{500}$)\\,$\\geq\\,1$ and log\\,(F$_{160}$/F$_{350}$)\\,$\\geq\\,1.6$ to identify an additional 31 embedded massive star candidates with no associated star-formation tracers. Using this result we can build a picture of the present day star-formation of the complex, and by extrapolating an initial mass function, suggest a current population of $\\approx\\,2\\,\\times\\,10^{4}$ young stellar objects (YSOs) present, corresponding to a star formation rate (SFR) of 0.01\\,-\\,0.02 M$_{\\odot}$\\,\\,yr$^{-1}$. Comparing this resolved star formation rate, to extragalactic star formation rate tracers (based on the Kennicutt-Schmidt relation), we find the star formation activity is underestimated by a factor of $\\geq$\\,2 in comparison to the SFR derived from the YSO population. ",
        "introduction": "\\begin{figure*} \\vspace{-14pt} \\begin{center} \\includegraphics[width=1.1\\textwidth]{G305_Publish6.eps} \\vspace{-15pt} \\caption{Three-colour (blue\\,=\\,70\\,$\\mu$m, green\\,=\\,160\\,$\\mu$m, red\\,=\\,350\\,$\\mu$m) \\textit{Herschel} Hi-GAL image of the giant HII region G305. The 70\\,\\micron\\,\\,emission corresponds to hot dust emission ($\\geq\\,40$\\,K), while the 350\\,\\micron\\,\\,emission emanates from more cold dust ($\\leq\\,20$\\,K). The positions of both Danks 1 \\& 2 (circles), and WR 48a (cross) are overplotted.} \\vspace{-10pt} \\label{Figure:G305} \\end{center} \\end{figure*}  The G305 star-forming complex is one of the most massive and luminous star-forming regions in the Galaxy \\citep{Clark:2004vd, Hindson:2010jt}, centred on the two optically visible open clusters Danks 1 \\& 2 \\citep{Danks:1984xy}, and the Wolf-Rayet star WR 48a. Located inside the Scutum-Crux arm within the Galactic plane at $l$\\,=\\,305$^{\\circ}$, $b$\\,=\\,0$^{\\circ}$, and found at a distance of $\\sim$\\,4\\,kpc, it has a projected diameter of $\\sim$\\,30\\,pc, and an estimated age of some 3\\,-\\,5 Myr \\citep{Clark:2004vd}. Coincident with G305 are numerous signposts of massive star formation, some of which are located on the periphery of the central cavity, such as infrared hotspots, compact and ultra-compact (UC) HII regions, H$_{2}$O, OH, and methanol masers \\citep{Urquhart:2007sc, Hindson:2010jt}.  We define a massive star as a star with sufficient mass to produce detectable radio emission associated with a HII region (i.e. M\\,$>$\\,8\\,M$_{\\odot}$). They play a key role in the Universe; their presence has a profound effect on the stellar and planetary formation process, while also on the physical, chemical, and morphological structure of galaxies \\citep{Kennicutt:2005ni}. Their UV radiation output ionises the surrounding interstellar neighbourhood providing the principle source of heating in the interstellar medium (ISM), while also having an impact on subsequent star formation. Through various process, such as stellar winds and supernovae, their mechanical energy output serves as the energy and momentum inputs to the surrounding ISM, helping to sculpt the structure and energetics of the ISM, and hence in host galaxies \\citep{Zinnecker:2007mk}.  Though they are important, in the shaping and evolution of their host galaxies, the physics of the formation and evolution of massive stars is unclear. Firstly, there is the issue of high dust extinction making direct observation difficult; sites of massive star formation tend to be embedded within very opaque cloud cores, which suffer from visual extinction greater than 100 magnitudes \\citep{Bally:2005zf}. Related to this also, is the fact that massive stars are predominantly found to form within dense, stellar clusters \\citep{de-Wit:2005fk}. In all, this makes for a particularly difficult challenge in being able to study the origins of massive stars.  Fig.\\ref{Figure:G305} underlines the dynamic morphology of the region, with the central location  Danks 1 \\& 2 and WR 48a thought to be the main source of energy input and the driving force behind the expansion and clearing of the central diffuse HII region in the complex \\citep{Clark:2004vd}. The suggestion is an interaction between the centrally embedded sources and the surrounding cloud, with an occurrence of ongoing massive star formation being located in the hot dust emission sites (seen in blue in Fig.\\ref{Figure:G305}) on the periphery of the central cavity. With the presence of numerous and different epochs of star formation in one location, and the relative close proximity, the G305 star-forming complex affords us an exceptional opportunity to study the nature of massive star formation, and investigate the environmental impact this may have on the formation of future generations of stars \\citep{Elmegreen:2002dg, Elmegreen:1977bs}.  This paper follows in a series of previous studies of massive star formation within the G305 star-forming complex. Previously, \\cite{Hindson:2010jt} have focused on the reservoir for star formation within the region through observations of NH$_{3}$ emission, while tracing sites of active star formation through H$_{2}$O maser emission. To complement this, \\cite{Davies:2011gl} have recently carried out a near-infrared study of the two central open clusters Danks 1 and 2; while current work has focused on identifying compact radio emission that is indicative of UC HII regions towards G305 \\citep{Hindson:2012lr}.  In this paper, we present a far-infrared (far-IR) study of the G305 complex using \\textit{Herschel} \\citep{Pilbratt:2010eg}, in conjunction with radio continuum, H$_{2}$O maser, methanol maser, MIPS, and Red MSX Source survey data, with the aim of identifying sites of embedded massive star formation. In this study we are able to resolve the embedded population within G305, since the \\textit{Herschel} far-IR observations are unaffected by dust, and combined with mid-IR data, constrain the luminosities of individual YSOs. By incorporating the luminosities of the embedded massive star-forming population with the initial mass function (IMF), we are able to determine the SFR of G305 and investigate the star formation history of the region. This resolved, Galactic SFR can then be compared to extragalactic SFR indicators to test whether the two regimes are consistent with one another, and identify where fundamental differences may lie \\citep{Heiderman:2010io, Lada:2010tb}.  The present study serves as an example of how \\textit{Herschel} data can be applied to Galactic star-forming regions, such as G305, in order to identify the high-mass stellar content of such complexes, and how the star formation activity can be inferred from this population. Following from this, a comprehensive YSO counting approach, similar to that conducted by \\cite{Povich:2011oj} for the Carina complex, will be conducted to tackle the incompletenesses present at the intermediate, and low-mass range in a following paper by the authors. By combining \\textit{Spitzer} GLIMPSE \\citep{Benjamin:2003qy}, and VISTA VVV \\citep{Minniti:2010fk} data alongside our present datasets we aim to conduct a complete census of the YSO population of G305, studying the physical parameters of these sources, their evolutionary stage, and spatial distribution within the complex.  This paper is organised into the following sections; in Section 2 we present and discuss the details of the observations and the data analysis, in Section 3 we present methods to identify sources detected in the far-IR through consideration of their spectral energy distribution, in Section 4 we present our discussion and place G305 in context with other star-forming regions, and in Section 5 we summarise our results. ",
        "conclusions": "In this section we discuss the global distribution of the embedded massive star-forming population within the G305 complex, and the general method of identification of this population; through the combination of the SED morphology and derived parameters. Using these identified sources, we are able to comment on the present-day star formation activity of G305, and place it into the context of Galactic star formation.  \\subsection{General Properties Of Sources Within G305}  We firstly consider the global properties of G305 obtained from SED fitting, by both modified blackbody and \\cite{Robitaille:2006mo} YSO models, for each association identified within the G305 region. For our fitting we consider sources with detections of N$_{data}$\\,$\\geq$\\,3 as acceptable to fit an SED to, since we have three free parameters in our modified blackbody fitting, sources with limited detections will unlikely produce a reliable SED; this leaves us with some 503 sources. From these fits, we deem sources with $\\chi^{2}$\\,/\\,N$_{data}\\,\\leq$\\,2 \\citep{Povich:2011oj}, as having a reliably fitted SED, yielding a total of 359 well fitted sources.  In Fig. \\ref{Figure:SED Parameter Histograms} we present the distribution of the modified blackbody parameters of the dust temperature, dust emissivity index, and the dust temperature, along with the bolometric luminosities obtained from \\cite{Robitaille:2006mo} model SED fitting. We find that dust temperatures lie within the range of 10 to 42\\,K, with a median value of $\\approx$\\,21\\,K. For the dust emissivity index, a range between 0.8 to 2.8 is found, with the median of $\\approx$\\,1.8. The bolometric luminosity is sampled between 10 to 10$^{4}$\\,\\lsol, with a median of  $\\approx$\\,300\\,\\lsol, while the dust mass lies between 1 to 10$^{4}$\\,\\msol, with a median of  $\\approx$\\,540\\,\\msol.  \\begin{figure} \\mbox{\\subfigure{\\includegraphics[width=0.23\\textwidth,angle=-90]{Temp.eps}} \\subfigure{\\includegraphics[width=0.23\\textwidth,angle=-90]{Beta.eps}}} \\mbox{\\subfigure{\\includegraphics[width=0.23\\textwidth,angle=-90]{L_Bol.eps}} \\subfigure{\\includegraphics[width=0.23\\textwidth,angle=-90]{Mass.eps}}} \\caption[dum]{Distribution of parameters for all sources with a reliable SED fit within G305 (solid black line), compared to the 16 identified candidate embedded massive star-forming regions (red dot-dashed line). \\textit{Top left}: dust temperature. \\textit{Top right}: dust emissivity index. \\textit{Bottom left}: bolometric luminosity. \\textit{Bottom right}: dust mass.} \\label{Figure:SED Parameter Histograms} \\end{figure}  \\subsection{A Far-IR Selection Criterion For Embedded Massive Star Formation}  \\begin{figure*} \\begin{center} \\includegraphics[width=0.65\\textwidth,angle=-90]{Embedded_MSF.eps} \\caption{Location of candidate embedded massive star-forming regions and their relative association overlaid onto a 160\\,$\\mu$m greyscale Hi-GAL image; ATCA radio data (blue circles), MMB methanol masers (cyan asterisks), RMS sources (red crosses), H$_{2}$O masers (purple diamonds). Numerous regions are found to be associated with multiple star formation tracers.} \\vspace{-10pt} \\label{Figure:G305 Embedded Massive Star-Forming Associations} \\end{center} \\end{figure*}  Previous studies that have attempted to distinguish a sample of MYSOs have relied on colour selection criteria based on IRAS data \\citep{Palla:1991kx,Molinari:1996uq,Walsh:1997sw,Sridharan:2002fj}, however these are restricted by selection effects, limitations in sensitivity and resolution, and also their incomplete Galactic coverage. The majority of these studies are biased in that they rely on bright IRAS sources, that with a resolution of 5\\arcmin\\,\\,at 100\\,\\micron\\,, are often found removed from the dense, confused regions along the Galactic plane where the majority of star formation is expected to be found. The result of these studies is a restricted sample that, due to selection issues, may not entirely represent the general MYSO population, making any extension to study massive star-formation problematic \\citep{Urquhart:2008xr}.  It has been suggested that both the SED morphology, and the bolometric luminosity can prove effective in determining a sample of embedded massive star-forming objects \\citep{Molinari:2008qy}. With the bulk of the YSO emission being in the far-IR portion of the SED, due to cold dust \\citep{Molinari:2008tj}, the \\textit{Herschel} Hi-GAL observations are ideal in accurately constraining the luminosity.  From the candidate associations identified, those sources found to have either a radio, MMB, H$_{2}$O maser, or RMS counterpart are known sites of massive star formation; tending to be luminous sources (i.e. $>$\\,10$^{3}$\\,L$_{\\odot}$). Using this sub-sample we are able to identify a sample of known embedded massive star-forming regions from the physical properties derived from, and morphology of, their SEDs (shown to peak at 100\\,\\micron\\,, as shown in Fig. \\ref{Figure:Robitaille SED Fits}). We refer to this population of Hi-GAL sources that are associated with radio, MMB, H$_{2}$O maser or RMS counterparts as the embedded population. Added to this, we employ a selection cut in the bolometric luminosity of 10$^{3}$\\,\\lsol\\,\\,, that corresponds to the minimum spectral type that we define as a massive star (i.e. M\\,$>$\\,8\\,M$_{\\odot}$). In total we find some 16 candidate embedded massive star-forming regions, that match these selection criteria; their respective properties and relative associations are shown in Table \\ref{tab:Embedded Massive Star-Forming Regions Parameters}. We also note the location of these associations, shown in Fig. \\ref{Figure:G305 Embedded Massive Star-Forming Associations}, with the majority being located in the hot dust locations along the periphery of the central cavity, suggesting an interaction between the central sources of Danks 1 \\& 2, and the surrounding material.  Based on this sample of 16 embedded massive star-forming regions, we are able to derive a two-colour selection criterion to identify the overall embedded population of the G305 complex. We find that the 70\\,-\\,500\\,\\micron\\,\\,and 160\\,-\\,350\\,\\micron\\,\\,colours are most sensitive to the embedded population, shown as asterisks in Fig. \\ref{Figure:FIR Colour Cut}. As can be seen from Fig. \\ref{Figure:FIR Colour Cut}, the embedded population are tightly confined to one area of the colour plot, and can be distinguished from the remaining G305 population, shown as circles. From this we suggest a far-IR colour selection criterion for embedded massive star-forming regions of log\\,(F$_{70}$/F$_{500}$)\\,$\\geq\\,1$ and log\\,(F$_{160}$/F$_{350}$)\\,$\\geq\\,1.6$, yielding a further 31 embedded massive star candidates with no associated emission and luminosities $>$\\,10$^{3}$\\,L$_{\\odot}$, as shown in Fig. \\ref{Figure:FIR Colour Cut}. From these 31 candidates we find the faintest source to have a peak 70\\,\\micron\\ flux of 1.02 Jy/pixel, and compared to the 90\\% recovery rate of sources discussed in Sect. 2.1, we do not expect any further more deeply-embedded massive star-forming regions to be found within G305.  Currently, the nature of these 31 candidate embedded massive star-forming regions is unclear; from Fig. \\ref{Figure:FIR Colour Cut} these candidates are predominately found with bolometric luminosities of $\\approx$\\,10$^{3}$\\,\\lsol\\, however at least 3 sources are found a luminosities $\\geq$\\,10$^{4}$\\,\\lsol\\ with no corresponding tracer of massive star formation. A lack of association to ATCA radio sources could be accounted for by localised noise found towards bright large-scale emission present with G305, confusing possible associations to real compact emission present (see \\cite{Hindson:2012lr} Sect. 2.2 for a detailed discussion of the ATCA data reduction process). Another possibility could be the strong variability of both methanol \\citep{Green:2012lr} and water masers \\citep{Breen:2010eq}, to such an extent that they display no common features at the present epoch. Aside from the possible reasons for lack of associated tracers, this sample of candidates may also suggest an earlier, very young embedded phase present within G305. The possible nature of these additional candidates is particularly interesting, and warrants further investigation at a later date.  \\begin{table} \\scriptsize \\addtolength{\\tabcolsep}{-4pt} \\caption{Derived physical properties for all identified embedded massive star-forming regions, from both modified blackbody fits and \\cite{Robitaille:2007nv} SED fitting technique, along with found associations.} \\begin{centering} \\begin{tabular}{lcccccc} \\hline Hi-GAL & $\\beta$ & T & $\\tau_{500\\mu m}$ & M$_{Dust}$ & L$_{Bol}$ & Association(s)\\\\ Source Index & & (K) & (10$^{-3}$) & (\\msol) & (10$^{3}$\\,\\lsol) & \\\\ \\hline 938   & 1.0 & 35 & 20.1 & 4300 & 48.6 &  MMB, H$_{2}$O maser\\\\ 945   & 1.1 & 29 & 15.4 & 3900 & 7.95 & RMS \\\\ 972   & 1.2 & 41 & 11.4 & 1200 & 20.8 & ATCA, MMB \\\\ 1184 & 1.1 & 25 & 27.1 & 4700 & 11.1 & MMB \\\\ 1800 & 1.4 & 35 & 29.4 & 3500 & 60.8 & ATCA, MMB, RMS \\\\ 1804 & 1.0 & 32 & 24.8 & 5200 & 17.4 & ATCA \\\\ 2114 & 1.7 & 24 & 1.29 & 200   & 2.91 & RMS \\\\ 2153 & 1.2 & 29 & 9.14 & 3700 & 4.13 & MMB, RMS, MIPS \\\\ 2212 & 1.1 & 23 & 14.3 & 4400 & 4.68 & RMS, MIPS \\\\ 2363 & 1.1 & 27 & 14.2 & 3200 & 12.2 & ATCA, RMS \\\\ 2383 & 1.0 & 36 & 20.0 & 1700 & 57.7 & ATCA, MMB, RMS \\\\ 2627 & 1.4 & 28 & 4.18 & 1500 & 3.65 & MMB \\\\ 2902 & 1.0 & 37 & 42.2 & 4200 & 48.6 & MMB\\\\ 2923 & 1.1 & 24 & 16.1 & 3700 & 3.24 & MMB \\\\ 2994 & 1.0 & 29 & 23.1 & 3900 & 7.95 & MMB, MIPS \\\\ 3032 & 1.4 & 28 & 4.57 & 1500 & 3.75 & MMB, RMS\\\\ \\hline \\end{tabular} \\label{tab:Embedded Massive Star-Forming Regions Parameters} \\end{centering} \\end{table}  \\subsection{The Present-day Star Formation Rate Of The G305 Complex}  With a sample of embedded massive star-forming objects, one opportunity that is available is to study the recent star-forming history of the complex. Determining the star formation rate (SFR) at a local level is crucial in determining the global Galactic SFR, helping to unveil any mechanisms that may lead to global scaling laws \\citep{Molinari:2010gl}. The SFR, along with the IMF, express the population of massive stars within the Galaxy, and determine what the impact on the local environment is, such as the composition of the ISM, the rate of feedback from massive stars, and the rate of conversion of gas into stars \\citep{Calzetti:2010aa}.  Given our position within the Galactic disk, direct SFR indicators using optical/UV tracers will fail to reproduce an accurate SFR due to high extinction rates of the dusty ISM. However, far-IR observations, unaffected by extinction, provide us with the ability to resolve the YSO population associated with HII regions, allowing constraints on the IMF and stellar ages, yielding a detailed star formation rate of Galactic HII regions \\citep{Chomiuk:2011fk}.  A SFR derived from a resolved, YSO counting approach, or from that inferred from the total infrared luminosity has to assume a `steady-state' approximation to reliably trace the star formation activity of the region. The assumption in these calculations is that both the rate of massive star formation, and the rate that massive stars evolve off the main sequence, is in approximate equilibrium \\citep{Krumholz:2012lr}. For this to be true, the requirement is that the age of the region be longer than that of the UV emitting population used to trace the SFR \\citep{Kennicutt:1998fk}. We show below that this is the case for a realistic star-forming timescale.  \\begin{figure} \\begin{center} \\includegraphics[width=0.4\\textwidth,angle=-90]{FIR_Cut_70-500.eps} \\includegraphics[width=0.4\\textwidth,angle=-90]{FIR_Cut_160-350.eps} \\caption{Colour - \u00adcolour plots of all Hi-\u00adGAL sources found in the G305 field (blue circles), and known sites of massive star formation (red asterisks). Dashed lines indicates the boundary of the region used for distinguishing sites of embedded massive star formation from other sources, at a luminosity of $>$\\,10$^{3}$\\,L$_{\\odot}$.} \\label{Figure:FIR Colour Cut} \\end{center} \\end{figure}  \\subsubsection{The Embedded Massive Star Formation Rate}  With our sample of 16 identified embedded massive star-forming regions, and the further 31 embedded massive star candidates found, we are able to comment on the present star formation history of the G305 complex. If we make the assumption that for each embedded massive star-forming region identified, the most massive star present produces the majority of the bolometric luminosity and is also accompanied by a cluster of lower mass stars, we are able to scale the IMF accordingly; for our calculations, we adopt a simple IMF proposed by \\cite{Salpeter:1955xv}. By comparing the calculated bolometric luminosity for each candidate region found, to the luminosities calculated for a sample of MYSOs from \\cite{Mottram:2011fs}, we are then able to estimate the most massive star for each region.  In order to extrapolate the IMF from observed mass to lower mass, we are obliged to select both a lower and upper mass limit that all observed YSOs fall within the selected range. For our purposes we adopt a lower mass of 0.1\\,\\msol\\,\\,and an upper limit of 50\\,\\msol, as used when calculating the Galactic SFR \\citep{Robitaille:2010eq}.  By adopting a Salpeter IMF, which best samples the high-mass tail of the IMF \\citep{Zinnecker:2007mk}, and assuming a constant SFR therefore, we arrive at $\\approx$\\,10$^{4}$ YSOs present, which corresponds to a total mass in stars of $\\approx\\,8\\,\\times$\\,10$^{3}$ M$_{\\odot}$. Since we consider the total mass in stars, we select a typical timescale for that mass to assemble, which would be the time to reach the pre main-sequence of 0.5\\,Myr \\citep{Offner:2011ph}, from this we attain a SFR for G305 of $\\approx$\\,0.01\\,-\\,0.02 M$_{\\odot}$\\,\\,yr$^{-1}$. In this scenario, the `steady-state' approximation should hold since the timescales of both MYSOs, $\\approx$\\,10$^{4}$\\,yr \\citep{Mottram:2011fs}, and UC HII regions, $\\approx$\\,10$^{5}$\\,yr \\citep{Comeron:1996fr}, are some 5\\,-\\,50 times shorter than our assumed timescale for the star formation within G305 of $\\approx$\\,0.5\\,Myr.  A similar approach has been taken in \\cite{Hindson:2012lr}, using the identified UC HII population (some five identified in total) of G305 to derive a SFR, over the last 0.5\\,Myr, of 0.002\\,-\\,0.004 M$_{\\odot}$\\,\\,yr$^{-1}$. This rate is considered a lower limit due to the incompletenesses in the ATCA radio data, with a \\textit{uv} cut implemented to emphasise the compact, small scale radio emission associated with UC HII regions. For comparison, \\cite{Davies:2011gl}, have determined the SFR for the G305 complex over the last 5\\,Myr from the calculated age and mass of the two central open clusters Danks 1 \\& 2. The SFR was found to be $\\approx$\\,0.002\\,-\\,0.005 M$_{\\odot}$\\,\\,yr$^{-1}$, which is comparable to our derived SFR from counting the embedded YSO population. However, taking our derived SFR of $\\approx$\\,0.01\\,-\\,0.02 M$_{\\odot}$\\,\\,yr$^{-1}$, it is clear that the star formation activity of G305 has not remained constant over the last 5\\,Myr. If this were to be the case, then we would expect to observe some 75\\,000 M$_{\\odot}$ of stars to be observed within the complex, which is entirely not the case. The `collect and collapse' model of star formation, proposed by \\cite{Elmegreen:1977bs}, requires a period of time after the formation of the central ionising source(s) for material within the surrounding molecular clouds to be swept-up by the expansion of the HII region. If this expansion continues for a sufficient time, the swept-up shell of material becomes self-gravitating and is then expected to fragment, entering a phase of collapse possibly leading to the formation of new stars. Following the approach taken by \\cite{Dale:2007nu}, who have calculated this `fragmentation time-scale' for a uniform molecular cloud of pure molecular hydrogen, we take the ionising photon flux of the G305 complex from \\cite{Clark:2004vd} to estimate a similar timescale. We find that after the formation of Danks 1 \\& 2, there would be a delay of $\\approx$\\,2.4\\,Myr until the next generation of star formation occurred within the complex. This would support the scenario that the star formation within G305 was not continuous, but more likely characterised by punctuated star formation over the lifetime of the complex.  We find also that our derived SFR for G305 is comparable to other well known massive star-forming complexes in the Galaxy, namely the Carina complex \\citep{Povich:2011oj}, and M17 \\citep{Chomiuk:2011fk}. We stress that the derived SFR value is based on a small sample of high mass stars, and has been extrapolated over a large range of stellar masses; when considering the lower mass stars present, their lifetimes may well be 1-2 orders of magnitude longer.  For completeness, the Galactic SFR is found to be $\\approx$\\,2 M$_{\\odot}$\\,\\,yr$^{-1}$ \\citep{Chomiuk:2011fk, Davies:2011af}, suggesting that a few tens to hundreds of G305 complexes are analogous to the entire star formation rate of the Milky Way. Using results from the Wilkinson Microwave Anisotropy Probe (WMAP), \\cite{Murray:2010fk} identify some 31 Galactic HII regions with an ionising flux greater than that of G305, of which some 18  WMAP sources constitute over half the total Galactic  ionising flux. Just as the IMF is dominated by the more massive stars present, the Galactic SFR is probably dominated by the few rigorous star-forming regions present.  \\subsection{Alternative Star Formation Rate Indicators Within G305}  Clearly the value for the SFR derived from the population of embedded massive star-forming regions identified will be an upper limit, since we have assumed a power-law slope (i.e. Salpeter IMF), and by extrapolating the IMF over a small sample of massive stars have overestimated the total mass in stars. To this must also be added issues of completeness, and assuming a timescale of $\\approx\\,10^{5}$\\,yrs  will be unrepresentative of the intermediate to low-mass YSOs present. However the value should prove a good assumption for the upper limit, that can then be compared to other SFR tracers that are independent of the resolved massive stellar population within G305.  What is key here is that we are able to resolve the YSO population within Galactic HII regions, such as G305, and use both the IMF and stellar timescales to derive a SFR. This can then be contrasted and calibrated to extragalactic emission tracers, such as the total IR luminosity, to determine whether Galactic SFRs are consistent with extragalactic SFR indicators. We next consider the SFR using tracers that are independent of the identified YSO population. Table \\ref{tab:SFR Indicators} lists the calculated SFRs for G305 using numerous tracers, with reference to each approach.  \\subsubsection{The Relation Between Star Formation Rate and Molecular Cloud Mass}  Recent work by \\cite{Heiderman:2010io} and \\cite{Lada:2010tb} on the star formation activity of molecular clouds within 0.5\\,kpc of the Sun, suggest that the star formation rate scales linearly with the molecular cloud mass. \\cite{Lada:1992fk} showed that active star formation is to be found primarily in high volume density regions of molecular clouds, with star formation favouring very massive, dense cores. The expectation for the star formation activity is thus that there is a tight correlation with the amount of high extinction material present in molecular clouds.  \\cite{Lada:2010tb} illustrate for a sample of local molecular clouds, that by comparison of the cumulative mass to YSO content as a function of extinction (see \\cite{Lada:2010tb} Fig. 3), a marked minimum dispersion of the cumulative mass is found at A$_{V}$\\,=\\,7.3\\,$\\pm$\\,1.8 mag. The proposition is that above this minimum, the cloud mass is directly related to the star formation activity and hence the SFR within the clouds. It is shown that this high extinction value corresponds to an equally high volume density of n(H$_{2}$)\\,$\\approx$\\,10$^{4}$\\,cm$^{-3}$ \\citep{Lada:2010tb}. There is also evidence this linear relation holds for extragalactic molecular clouds. A tight correlation between the total IR luminosity and the luminosity of the HCN molecule, which itself requires high densities ($\\textgreater$\\,10$^{4}$\\,cm$^{-3}$) to be excited to a detectable level, has been found for both Galactic cores \\citep{Wu:2005fk}, and for a sample of normal spirals and starburst galaxies \\citep{Gao:2004fk}. These studies suggest that the linear relation holds for dense interstellar gas both on a Galactic and extragalactic scale, underlying a physical relation that links star formation and galaxy evolution.  From this linear correlation, a SFR is derived of the form:  \\begin{equation} SFR=4.6\\,\\pm\\,2.6\\,\\times\\,10^{-8}\\,M_{0.8}\\,\\,\\rm{[M_{\\odot}\\,\\,yr^{-1}]} \\label{eq:Dense gas SFR equation} \\end{equation} where M$_{0.8}$ corresponds to the cloud mass, in solar masses, above an extinction threshold of A$_{K}$\\,$\\approx$\\,0.8 mag, which is derived from the visual extinction of A$_{V}$\\,=\\,7.3\\,$\\pm$\\,1.8 mag mentioned earlier.  For each IR source identified by \\textit{Herschel} across the G305 complex we determine the physical radius at 250\\,\\micron, as this wavelength offers the optimum combination of signal-to-noise, and angular resolution. To estimate the physical diameter for each source, we firstly deconvolve the source size from the Gaussian beam \\citep{Thompson:2004qy}:  \\begin{equation} \\Theta^{2}_{\\rm{Source}}=\\Theta^{2}_{\\rm{Obs}} - \\Theta^{2}_{\\rm{Beam}} \\label{eq:Convolution Equation} \\end{equation} where $\\Theta^{2}_{Obs}$ is the source size estimated from the FWHM of the Gaussian fitting, and $\\Theta^{2}_{Beam}$ is the \\textit{Herschel} beam size at 250\\,\\micron. If we then place each source at a distance of 4\\,kpc, and assume spherical geometry and a uniform density, we can determine those sources found to be above the critical density threshold. From this we are able to determine the mass of the dense gas within G305, then using the \\cite{Lada:2010tb} assumption that dense gas is associated with the star formation activity, we find a dense gas mass of $\\approx$\\,3\\,$\\times$\\,10$^{5}$\\,M$_{\\odot}$; for comparison, the total molecular mass traced by NH$_{3}$ is found to be $\\approx$\\,6\\,$\\times$\\,10$^{5}$\\,M$_{\\odot}$ \\citep{Hindson:2010jt}. We note that this mass is an approximation, since it is unlikely that the density of each source is uniform.  By combining this mass with Equation (\\ref{eq:Dense gas SFR equation}), we obtain a dense gas derived SFR of 0.006\\,-\\,0.02 M$_{\\odot}$\\,\\,yr$^{-1}$. The result is found to be in good agreement with the embedded massive star formation rate derived earlier, and goes some way toward confirming the \\cite{Lada:2010tb} assumption by extension to the star formation activity of embedded massive star-forming regions.  \\subsubsection{The 70\\,\\micron\\ Emission Star Formation Rate}  In contrast to deriving a value of the SFR from all identified embedded massive star-forming regions, we can also approach measuring the SFR by considering the total infrared flux (TIR) of the giant G305 HII region. On an extra-galactic perspective current SFRs are calculated using a tracer of UV photon emission from YSOs, and also spectral synthesis models \\citep{Kennicutt:1998fk}. In this case, observations of HII regions prove ideal measures of current star formation. However, some fraction of the total UV emission will be obscured by the presence of dust, thus bolometric IR observations of dust (i.e. TIR) will provide an excellent means to recover the extinguished UV photon emission, with the dust absorption highly peaked in the UV and re-emission being in a broad spectral range of mid-to-far-IR \\citep{Kennicutt:1998fk}. The conclusion from this is that the TIR will provide the best indicator of SFR obscured by the presence of dust.  Observationally, an advantage would be the use of a single-band star formation indicator, such as UV, H$\\alpha$, 8\\,\\micron, 24\\,\\micron\\,\\,etc., however each of these has its own complications. In the case of UV and optical lines, corrections need to be taken into consideration due to large extinction. Whereas 8 and 24\\,\\micron\\,\\,emission strongly depends on the local environment, since the abundances of small dust grains that contribute to their emission depend greatly on metallicity and the presence of ionising radiation.  \\cite{Calzetti:2007xx} and \\cite{Dale:2005sz} note that 8\\,\\micron\\,\\,emission makes for an inaccurate SFR indicator since there is a large degree of variability of emission in galaxies with respects to SED shape and metallicity.  This strong variability at 8\\,\\micron\\,\\,is emphasised in Fig. \\ref{Figure:Robitaille SED Fits}, where variations of an order of magnitude exist between the best fit and good fit SED models. This variation can be accounted for by the disk inclination to the line of sight for the centrally embedded object, where the observed flux from a pole-on view can be 2\\,-\\,4 times greater than a more edge-on viewing angle \\citep{Whitney:2003zr}.  \\cite{Calzetti:2005oh} find that the SFR calculated from 24\\,\\micron\\,\\,emission itself varies strongly from galaxy to galaxy. On a local scale, the ratio of the 24\\,\\micron\\,\\,luminosity to SFR is found to be a reasonably accurate tracer, however when applied to other systems, such as starbursts and ultraluminous infrared galaxies (ULIRGs), the ratio is systematically higher. This variation in the 24\\,\\micron\\,\\,is found to be a factor of a few with respects to the observed SEDs, and may be grounded in the strong dependence on local galactic conditions; with ionising stars heating dust to different averaged temperatures, the 24\\,\\micron\\,\\,emission will be most sensitive to this. However, \\cite{Dale:2005sz} find that 70\\,\\micron\\,\\,emission may be an accurate monochromatic star formation indicator, since the 70-to-160\\,\\micron\\,\\,ratio is found to correlate well with local SFRs.  Recent work by \\cite{Lawton:2010rj} have determined an accurate monochromatic IR band that best approximates the obscured SFR in the Large Magellanic (LMC) and Small Magellanic Clouds (SMC), through IR aperture photometry of 16 LMC and 16 SMC HII regions, using \\textit{Spitzer} IRAC (3.6, 4.5, 8\\,\\micron) and MIPS (24, 70, 160\\,\\micron) bands. It is found from the IR SEDs of each HII region, that the majority peak at around 70\\,\\micron\\,\\,at all radii (10\\,-\\,400\\,pc) from the centrally ionising sources, and that the 70\\,\\micron\\,\\,emission most closely traces the size of each HII region as found using the TIR. The conclusion from this is that the 70\\,\\micron\\,\\,emission is the most likely suitable IR band to utilise as a monochromatic SFR indicator.  It has been argued by \\cite{Kennicutt:1998fk} that the TIR is the best obscured SFR indicator available for starburst galaxies. However, dust obscured star formation in HII regions are found to behave similarly, in that their environments are both very dusty and are sites of recent star formation. The \\cite{Kennicutt:1998fk} obscured SFR equation is of the form:  \\begin{equation} SFR=4.5\\,\\times\\,10^{-44}\\,L_{\\rm{TIR}}\\,\\,\\rm{[M_{\\odot}\\,\\,yr^{-1}]} \\label{eq:Kennicutt SFR Equation} \\end{equation} where L$_{TIR}$ is the TIR luminosity in erg\\,\\,s$^{-1}$, and the value 4.5\\,$\\times$\\,10$^{-44}$ is a constant derived from synthesis models, with assumptions on the IMF and star formation timescales \\citep{Kennicutt:1998fk}.  The TIR luminosity in Equation (\\ref{eq:Kennicutt SFR Equation}) can be substituted with the averaged 70\\micron\\,\\, luminosity, normalised by the TIR, while also applying an IR band specific constant. The monochromatic obscured SFR equation of \\cite{Lawton:2010rj} is found to be:  \\begin{equation} SFR=9.7(0.7)\\,\\times\\,10^{-44}\\,L(\\lambda)\\,\\,\\rm{[M_{\\odot}\\,\\,yr^{-1}]} \\label{eq:Monochromatic obscured SFR} \\end{equation} where L($\\lambda)$ is the observed luminosity in erg\\,\\,s$^{-1}$.  By employing aperture photometry of the whole G305 region, we obtain the cumulative 70\\,\\micron\\,\\, flux, f$_{\\nu}(\\lambda)$, and from this are able to calculate the monochromatic luminosity at 70\\,\\micron\\,\\,\\citep{Calzetti:2010zh}:  \\begin{equation} L(\\lambda)=4\\,\\pi\\,d^{2}\\,\\left(\\frac{c}{\\lambda}\\right)\\,f_{\\nu}(\\lambda)\\,\\,\\rm{[erg\\,\\,s^{-1}]} \\label{eq:Monochromatic luminosity} \\end{equation} where d is the distance to the G305 complex, in m.  Using the value found for the observed luminosity at 70\\,\\micron\\,, with Equation (\\ref{eq:Monochromatic obscured SFR}), we obtain an obscured SFR for G305 of 0.002\\,-\\,0.005 M$_{\\odot}$\\,\\,yr$^{-1}$. A similar approach has been suggested by \\cite{Li:2010hr}, who also determine a monochromatic SFR indicator at 70\\,\\micron\\,, yet calibrate their SFR tracer not with the TIR luminosity but rather with the combined H$\\alpha$, and 24\\,\\micron\\,\\,luminosity. For completeness, using the \\cite{Li:2010hr} tracer, we derive a SFR of 0.004\\,-\\,0.008 M$_{\\odot}$\\,\\,yr$^{-1}$, which is in approximate agreement with that using the \\cite{Lawton:2010rj} approach.  We can directly compare these results to that derived from the total Lyman continuum photon rate of G305, where we find an SFR of 0.002\\,-\\,0.004\\,M$_{\\odot}$\\,\\,yr$^{-1}$ (Hindson et al. in prep). We note that both these two independent tracers are in excellent agreement, however are found to a factor of $\\geq$\\,2 lower than that derived from the embedded massive star-forming population. A similar result is found by \\cite{Chomiuk:2011fk}, who find that the SFR for M17 estimated from both the Lyman continuum and 24\\,\\micron\\,\\,emission is underestimated by a factor of $\\geq$\\,2 in comparison to the SFR derived from YSO counting.  This discrepancy may simply be that there are elementary differences in the measurements between Galactic and extragalactic observations. \\cite{Lawton:2010rj} note that Equation (\\ref{eq:Monochromatic obscured SFR}) extends to HII regions measured at projected distances of 52\\,kpc and 61\\,kpc for the LMC and SMC respectively. With G305 some $\\sim$\\,4\\,kpc distant, the relation between the SFR and luminosity at 70\\,\\micron\\,\\,may indeed break down, with the effects of individual protostars becoming more important, due to the larger spatial resolution. It may also be the case that the `steady-state' assumption breaks down in this comparison. Though it is normally true that the lifetime of the region observed is longer than for the individual objects for extragalactic realms, this may not hold for observations of Galactic regions that tend to be smaller, and with shorter dynamical timescales.  \\subsubsection{A Galactic - Extragalactic Comparison}  By comparing the derived SFR from numerous tracers, shown in Table \\ref{tab:SFR Indicators}, what is immediately apparent is the disparity between the rates derived from the resolved stellar population and those from extragalactic tracers; there is lack of consistency between the two, with extragalactic tracers tending to underestimate the SFR derived from resolved Galactic SFRs. This circumstance between the two SFR regimes has been noted by several authors \\citep{Heiderman:2010io,Lada:2010tb,Chomiuk:2011fk}, where there appears a distinct underestimation for Galactic H II regions. \\cite{Heiderman:2010io} make the suggestion that this difference may be accounted for by the inclusion of diffuse gas, in the standard Kennicutt-Schmidt relation, that is below the critical density threshold for star formation, as suggested by \\cite{Lada:2010tb}. Since extinction maps are not readily available to determine the surface density of gas in extragalactic studies, CO maps are often employed instead. \\cite{Heiderman:2010io} find that using CO as a gas tracer, for a sample of local molecular clouds, leads to an underestimate in mass of  $\\gtrsim$\\,30\\,\\% compared to that obtained using extinction maps. The result of this would essentially push down the estimated SFR from the \\cite{Kennicutt:1998fk} relation, and may go some way in accounting for the dissimilarity between extragalactic regions and more local, Galactic ones.  However, recent work by \\cite{Krumholz:2012lr} seems to suggest a unified star formation law, with objects ranging from both low mass Solar neighbourhood clouds through to sub-mm galaxies in agreement with one distinct star formation law. What is advocated in this law is that the SFR, within a variety of scales, is simply $\\approx$\\,1\\,$\\%$ of the molecular gas mass per local free-fall time. This volumetric approach suggests a local, universal star formation law that is applicable to Galactic and extragalactic observations (see \\cite{Krumholz:2012lr} Fig. 3), bridging the gap between the apparent disparity in the two regimes. This law is affected solely by local variations, such as the gas condition, with more global Galactic/galaxy-scale properties, such as the orbital period, having no impact on the SFR in so far as they do not change the local properties of star-forming regions.  Conversely, \\cite{Lada:2011lr} conclude that a universal star formation law that is applicable from the Milky Way through to near-IR selected (BzK) galaxies, is simply directed by the amount of dense molecular gas that can accumulate within a star-forming region. In the majority of situations, only 10\\,$\\%$ of the total mass within a molecular cloud is at a sufficient density, n(H$_{2}$)\\,$\\geq$\\,10$^{4}$\\,cm$^{-3}$ \\citep{Lada:2010tb}, to actively form stars. Clearly there is a disparity between the two proposed universal SFR laws; \\cite{Krumholz:2012lr} favouring gas surface densities and local free-fall times as crucial, while \\cite{Lada:2011lr} advocate gas surface densities and the fraction of dense gas as the pivotal factors. It therefore seems that more work is needed to describe the underlying nature of a universal star formation law, if such a law is to be found.  What is apparent, when measuring the Galactic SFR, is the need for an accurate means to compare the Milky Way to other galaxies, in order to allow us to extend the more detailed Galactic analysis to other systems and to test the discrepancy between the two regimes. Continued, multiwavelength analysis of Galactic HII regions, now including \\textit{Herschel} Hi-GAL, will in part aid with this. Through extended study across a wide range of star-forming regions, an accurate determination of the IMF, and with that the SFR, can be achieved, which scaled up from a more local level to a global, Galactic level, will allow for the consideration of how these crucial properties vary as a function of environment across the Milky Way (see \\cite{Veneziani:2012fk} for a detailed study of the Herschel Science Demonstration Phase fields). This should help in a better understanding of how the SFR can accurately be measured on both Galactic, and extragalactic scales, and lead to a more unified calibration.  \\begin{table} \\scriptsize \\addtolength{\\tabcolsep}{-5pt} \\caption{Calculated SFR for G305 using multiple SFR tracers.} \\begin{centering} \\begin{tabular}{lcc} \\hline SFR Tracer & SFR & Reference \\\\ & (M$_{\\odot}$\\,\\,yr$^{-1}$) & \\\\ \\hline Embedded Massive Stars & 0.01\\,-\\,0.02 & This Paper \\\\ Dense Gas & 0.006\\,-\\,0.02 & \\cite{Lada:2010tb} \\\\ & & \\\\ UC HII Regions & $\\geq$\\,0.002\\,-\\,0.004 & \\cite{Hindson:2012lr} \\\\ Danks 1 \\& 2  & 0.002\\,-\\,0.005 & \\cite{Davies:2011gl} \\\\ 70\\,\\micron\\,\\, Emission & 0.002\\,-\\,0.005 & \\cite{Lawton:2010rj} \\\\ & 0.004\\,-\\,0.008 & \\cite{Li:2010hr} \\\\ Lyman Continuum & 0.002\\,-\\,0.004 & Hindson et al. in prep.\\\\ \\hline \\end{tabular} \\label{tab:SFR Indicators} \\end{centering} \\end{table}"
    },
    "1208/1208.5834_arXiv.txt": {
        "abstract": "We present LITTLE THINGS (Local Irregulars That Trace Luminosity Extremes, The \\HI\\ Nearby Galaxy Survey) that is aimed at determining what drives star formation in dwarf galaxies. This is a multi-wavelength survey of 37 Dwarf Irregular and 4 Blue Compact Dwarf galaxies that is centered around \\HI-line data obtained with the National Radio Astronomy Observatory (NRAO) Very Large Array (VLA). The \\HI-line data are characterized by high sensitivity ($\\leq1.1$ mJy beam$^{-1}$ per channel), high spectral resolution ($\\leq$2.6 \\kms), and high angular resolution ($\\sim$6\\arcsec). The LITTLE THINGS sample contains dwarf galaxies that are relatively nearby ($\\leq$10.3 Mpc; 6\\arcsec\\ is $\\leq$300 pc), that were known to contain atomic hydrogen, the fuel for star formation, and that cover a large range in dwarf galactic properties. We describe our VLA data acquisition, calibration, and mapping procedures, as well as \\HI\\ map characteristics, and show channel maps, moment maps, velocity-flux profiles, and surface gas density profiles. In addition to the \\HI\\ data we have {\\it GALEX} UV and ground-based $UBV$ and \\ha\\ images for most of the galaxies, and $JHK$ images for  some. {\\it Spitzer} mid-IR images are available for many of the galaxies as well. These data sets are available on-line. ",
        "introduction": "\\label{sec-intro}  Dwarf Irregular (dIrr) galaxies are the most common type of galaxy in the local universe (see, for example, Dale \\et\\ 2009). They are also the closest analogs to the low mass dark matter halos that formed after the Big Bang. In the $\\Lambda$CDM model, it is in these entities that the first stars formed. Yet, we do not understand what drives star formation on galactic scales even in nearby dwarfs, the simplest, most pristine local environments (see, for example, Hunter 2008, Bigiel \\et\\ 2010). To remedy this situation we have assembled multi-wavelength data on a large sample of relatively normal, nearby gas-rich dwarf galaxies, tracing their stellar populations, gas content, dynamics, and star formation indicators. This project is called LITTLE THINGS (Local Irregulars That Trace Luminosity Extremes, The \\HI\\ Nearby Galaxy Survey), and it builds on the THINGS project, whose emphasis was on nearby spirals (Walter \\et\\  2008).  LITTLE THINGS\\footnote[17]{http://www.lowell.edu/users/dah/littlethings/index.html} was granted $\\sim$376 hours with the Very Large Array (VLA\\footnote[18]{ The VLA is a facility of the National Radio Astronomy Observatory (NRAO). The NRAO is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. These data were taken during the upgrade of the VLA to the Expanded VLA or EVLA. In this paper we refer to the instrument as the VLA, the retrofitted antennae as EVLA antennae, and non-retrofitted antennae as VLA antennae. This emphasizes the hybrid nature of the instrument and distinguishes it from the far more powerful Jansky VLA or JVLA it has become since 2012. }) in B, C, and D configurations to obtain \\HI-line data of 21 dIrr and Blue Compact Dwarf (BCD) galaxies. We added another 20 galaxies from the VLA archives for a total sample of 37 dIrr and 4 BCD galaxies covering a large range of dwarf galactic parameter space. Our VLA data emphasize deep, high spatial (6\\arcsec) and spectral ($\\le$2.6 km s$^{-1}$) resolution in order to reveal clouds, shells, and turbulence in the interstellar medium (ISM) that could be important for creating regions of higher density star-forming gas. We are combining the \\HI-line data with ultraviolet (UV), optical, and infrared (IR) data to address the following questions:  {\\it What regulates star formation in small, gas-rich galaxies?} Studies of the star formation laws in galaxies show that star formation is very inefficient in dwarf galaxies and in the outer disks of spirals (Bigiel \\et\\ 2010; Ficut-Vicas \\et, in preparation). Furthermore, the star formation efficiency decreases as the gas density decreases. How then do atomic hydrogen clouds form in the low gas density environments of dwarfs that ultimately lead to star-forming molecular clouds?  {\\it What is the relative importance of sequential triggering by previous generations of stars for star formation in dwarf galaxies?} One generation of stars can trigger the formation of the next by rearranging the gas through winds and supernova explosions (see, for example, \\\"Opik 1953, Gerola \\et\\ 1980, Comins 1983, Dopita \\et\\ 1985, Mori \\et 1997, van Dyk \\et\\ 1998), but how important is this process? Observations of dwarfs show a better correlation between the star formation rate and the $V$-band surface brightness, which emphasizes Gyr old stars, than with any other measure, including the average gas density (Hunter \\et\\ 1998). Does this relationship imply a significant causal relationship in which existing stars are important for triggering new star formation?  {\\it What is the relative importance of random turbulent compression cloud formation in dwarf galaxies?} Turbulence can account for a wide range of phenomena that are indirectly related to star formation: the power law luminosity functions of \\HII\\ regions, the power law mass functions of clouds and clusters (Mac Low \\& Klessen 2004), the shape of the power spectra of \\HI\\ (Stanimirovi\\'c \\et\\ 1999, Elmegreen \\et\\ 2001, Dickey \\et\\ 2001, Zhang \\et\\ 2012a) and H$\\alpha$ (Willett \\et\\ 2005), the log-normal distribution of the H$\\alpha$ intensity in individual pixels in H$\\alpha$ images of dwarf galaxies, (Hunter \\& Elmegreen 2004), the correlation between region size and the star formation time scale (Efremov \\& Elmegreen 1998), and hierarchical structures in the ISM (Bastian \\et\\ 2007). But how important is turbulence in the formation of star-forming clouds, particularly in the realm of sub-critical gas density?  {\\it What happens to the star formation process in the outer parts of disks?} We find that many dwarfs have exponential stellar surface brightness profiles with a break and a change in slope. These breaks are similar to what are seen in the outer parts of spirals (e.g., de Grijs \\et\\ 2001; Kregel \\et\\ 2002; Pohlen \\et\\ 2002; Kregel \\& van der Kruit 2004; Pohlen \\& Trujillo 2006; Erwin \\et\\ 2008; Guti\\'{e}rrez \\et\\ 2011) and disks at high redshift (P\\'erez 2004). The break could imply a change in the star formation process at the break radius, but what is really happening there?  Here we describe the LITTLE THINGS project and present the \\HI\\ data. Science papers that are part of this project include: 1) ``Velocity Dispersions and Star Formation Rates in the LITTLE THINGS Dwarf Galaxies (Cigan \\et, in preparation), 2) ``Star Formation Laws in LITTLE THINGS Dwarfs: The case of DDO 133 and DDO 168'' (Ficut-Vicas \\et, in preparation), 3) ``Deep Radio Continuum Imaging of the Dwarf Irregular Galaxy IC 10: Tracing Star Formation and Magnetic Fields'' (Heesen \\et\\ 2011), 4) ``Deep Radio Continuum Observations of the Nearby Dwarf Irregular Galaxy IC 10. I. Nature of the Radio Continuum Emission'' (Heesen \\et, in preparation), 5) ``Surface Brightness Profiles of Dwarf Galaxies. I. Profiles and Statistics'' (Herrmann \\et\\ 2012), 6) ``Surface Brightness Profiles of Dwarf Galaxies. II. Color Trends and Mass Profiles'' (Herrmann \\et, in preparation), 7) ``The Stellar and Gas Kinematics of the LITTLE THINGS Dwarf Irregular Galaxy NGC 1569'' (Johnson \\et\\ 2012), 8) ``The Stellar Kinematics of the LITTLE THINGS Dwarf Irregular Galaxies DDO 46 and DDO 168'' (Johnson \\et, in preparation), 9) ``Deep 6 cm Radio Continuum Observations of the Nearby Dwarf Irregular Galaxy IC 1613'' (Kitchener \\et, in preparation), 10) ``High-resolution Mass Models of Dwarf Galaxies from LITTLE THINGS'' (Oh \\et, in preparation), 11) ``Star Formation in IC 1613\" (Simpson \\et, in preparation), 12) ``Outside-in Shrinking of the Star-forming Disk of Dwarf Irregular Galaxies\" (Zhang \\et\\ 2012b), and 13) ``\\HI\\ Power Spectrum Analysis of Dwarf Irregular Galaxies'' (Zhang \\et\\ 2012a). Other science that is in progress includes, but is not limited to, studies of the following: 1) Star Formation Processes in Blue Compact Dwarfs -- Ashley, 2) Molecular Cloud Structure and Dust at Low Metallicity -- Cigan, 3) Star Formation Laws of the full sample of galaxies -- Ficut-Vicas, 4) Surface Brightness Profiles and their relationship to \\HI\\ gas and two dimensional images -- Herrmann, 5) The Role of Turbulence in Star Formation in Dwarfs -- Hunter, 6) A Deep 6 cm Radio Continuum Survey of LITTLE THINGS -- Kitchener, and 7) Porosity of the Interstellar Medium in Dwarfs -- Simpson. ",
        "conclusions": ""
    },
    "1208/1208.2413_arXiv.txt": {
        "abstract": "A new modeling way of describing the continuous absorption of electromagnetic (EM) radiation in dense partially ionized hydrogen plasma is tested in this work. It is shown that the obtained results give a possibility of calculating spectral absorption coefficients which characterize the relevant absorption processes in partially ionized hydrogen plasmas with electron densities $N_{e} \\sim 10^{19}cm^{-3}$ and temperatures $T \\approx 2\\cdot 10^{4}K$. A key parameter of the used procedure is determined empirically. The calculation method is applied to wavelength region $300nm < \\lambda < 500nm$. The presented results can be of interest for dense laboratory plasmas as well as for partially ionized layers of different stellar atmospheres. ",
        "introduction": "\\label{sec:in}  By now, direct methods of determination of various plasma characteristics, based on quantum or classical statistical mechanics, have been developed only for practically fully ionized plasmas \\cite{kob95,ada94,ich87,ebe76, kra86, rin88, for89, mih04}. In the case of dense partially ionized plasmas, where the density of neutral particles (atoms) is close to the density of positively charged particles (ions), such rigorous methods do not exist at present. Recently, in \\cite{sre10}, this problem was discussed in connection with the transport properties of dense partially ionized plasmas. As for their optical characteristics, it is enough to remind that adequate calculation methods exist only for weakly and moderately ionized plasmas with electron density $N_{e} \\lesssim 10^{17} cm^{-3}$. As it is well known, the influence of the neighborhood on an exited atom can be neglected in such plasmas, as for example in the Solar photosphere \\cite{mih78, mih07a}, or be treated as a small perturbation and described within the framework of a perturbation theory \\cite{kob95, dam10, oma07, sea94, dim96, tka97, rei03, min06, SCCS98}.  In this paper we will consider the continuous absorption of EM (electromagnetic) radiation in dense partially ionized plasma, with electron density $N_{e} \\sim 10^{19} cm^{-3}$, temperature $T \\approx 20 000K$ and atom density $N_{a}\\approx N_{e}$. Plasmas with similar parameters are of interest from both the laboratory \\cite{gav01, dya06} and the astrophysical aspect. Here we keep in mind the plasma of the inner layers of the solar atmosphere, as well as of partially ionized layers of other stellar atmospheres, for example the atmospheres of DA and DB white dwarfs with effective temperatures between $10 000 K$ and $20 000 K$ (see \\cite{koe97, kie04,fin97}).  Due to the exceptional simplicity of the hydrogen atom, this research is starting with the hydrogen case. Under the mentioned conditions the continuous absorption of EM radiation in hydrogen plasmas are determined by the following radiative processes: \\begin{equation} \\label{eq:ph} \\varepsilon _{\\lambda} + H^{*}(n,l) \\to H^{+} + e_{k'}, \\end{equation} \\begin{equation} \\label{eq:inbs} \\varepsilon _{\\lambda} + e_{k} + H^{+} \\to e_{k'} + H^{+}, \\end{equation} \\begin{equation} \\label{eq:eH} \\varepsilon_{\\lambda} + \\left\\{ \\begin{array}{l} e_{k} + H(1s), \\\\ H^{-}(1s^{2}), \\end{array} \\right. \\to e_{k'} + H(1s), \\end{equation} \\begin{equation} \\label{eq:HH} \\varepsilon_{\\lambda} + \\left\\{ \\begin{array}{l} H^{+} + H(1s), \\\\ H_{2}^{+}(1\\Sigma^{+}_{g}), \\end{array} \\right. \\to H(1s) + H^{+}, \\end{equation} where $\\epsilon_{\\lambda}$ is the energy of a photon with wavelength $\\lambda$, $n$ and $l$ - the principal and the orbital quantum number of hydrogen-atom excited states, $e_{k}$ and $e_{k'}$- free electrons with energies $E = \\hbar^{2}k^{2}/2m$ and $E' = \\hbar^{2}k'^{2}/2m$ respectively, $m$ the electron mass and $\\hbar$ - Plank's constant.  In the considered region of $N_{e}$ and $T$ the processes of electron-$H^{+}$ and electron-$H(1s)$ inverse bremsstrahlung, photo-detachment of the negative ion $H^{-}(1s^{2})$, $H(1s)$-$H^{+}$ absorption charge exchange and photo-dissociation of molecular ion  $H_{2}^{+}(1\\Sigma^{+}_{g})$, which are described by Eqs.~(\\ref{eq:inbs}), (\\ref{eq:eH}) and (\\ref{eq:HH}), can be treated in the same way as in previous papers \\cite{mih93,mih94}. Therefore, photo-ionization processes (\\ref{eq:ph}) will be in the focus of attention in the next section, and here let us note only that in this field methods obtained by extrapolation (with minor modifications) of the methods developed for weakly and moderately non-ideal plasmas \\cite{dya98,gav01,vit04} have been used for such processes up until now. Moreover, for determination of absorption coefficients characterizing their influence in the region of $N_{e} \\gtrsim 10^{19} cm^{-3}$ methods based on Cramer's approximation \\cite{dya06,mof08} have been used so far. So, developing a new modeling way of describing the continual absorption in dense partially ionized plasmas, which is the main aim of this paper, is still an actual task.  We take as the landmarks hydrogen plasmas which were experimentally studied in \\cite{vit04}: with $N_e = 6.5\\cdot10^{18}cm^{-3}$ and $1.5\\cdot10^{19}cm^{-3}$, and $T =1.8\\cdot10^{4} K$ and $2.3\\cdot10^{4}K$ respectively. The presented modeling way is tested within the optical range of photon wavelengths: $300 nm \\le \\lambda \\le 500 nm$. A key parameter of the used numerical procedure is determined empirically. It is shown that this procedure already allows for determination of the spectral absorption coefficients characterizing all the relevant absorption processes in dense partially ionized hydrogen plasmas, at least in the regions $5\\cdot 10^{18} cm^{-3} \\lesssim N_e \\lesssim 5 \\cdot 10^{19}cm^{-3}$ and $1.6 \\cdot 10^{4} K \\lesssim T \\lesssim 2.5\\cdot 10^{4} K$.  The material of this paper is distributed over Sections 2 and 3. In Sec. 2 the following is presented: the approximation of the cut-off Coulomb potential, together with the reasons for applying just this approximation in the considered regions of $N_{e}$ and $T$; the way of obtaining all the partial spectral absorption coefficients, as well as their corresponding final expressions. In Sec. 3 the results of the calculations of the partial and the total absorption coefficients are presented and discussed. ",
        "conclusions": "}  \\subsection{The characteristics of the cut-off Coulomb potential}  In this paper the approximation of cut-off Coulomb potential (\\ref{eq:U0}) is applied to modeling the spectral absorption coefficient obtained in \\cite{vit04} in two experiments with hydrogen plasmas, which are treated as a short and a long pulse respectively. In the first case (short pulse) plasma with $N_e= 1.5 \\cdot 10^{19} cm^{-3}$ and $T = 2.3\\cdot 10^{4} K$ was studied, while in the second case (long pulse) - it was plasma with $N_{e}= 6.5 \\cdot 10^{18} cm^{-3}$ and $T =1.8\\cdot 10^{4} K$. In the experiments described in \\cite{vit04} plasmas with electron densities up to $\\approx 10^{19} cm^{-3}$ were created by pulse discharge in quartz capillary. Diagnostics of the plasma was carried out on the basis of optical measurements (at $\\lambda = 632.8 nm$), taking into account radial inhomogeneity of the plasma column. The temperature profile is defined from independent measurements of brightness and transparency at different distances from the center of the capillary. A detailed study is performed just for the two above mentioned examples.  On the basis of Fig.\\ref{fig::2} and Eq.~(\\ref{eq:pci}) it was found that the cut-off radius $r_{c}$ is equal to : $44.964$ a.u. for the short pulse, and $55.052$ a.u. for the long one. For these values of $r_{c}$ the solutions of Eq.~(\\ref{eq:Shred}) correspond to the energies of the realized bound state, which are presented in Tab's 1 and 2 respectively. The corresponding partial photo-ionization cross-sections $\\sigma(\\lambda;n,l,\\epsilon_{n,l})$ are obtained by means of Eq.~(\\ref{eq:sig}) for $n$, $l$ and $\\epsilon_{n,l}$ given in  Tab's 1 and 2. The behavior of these cross-sections is illustrated in Fig.\\ref{fig::3} and Fig.\\ref{fig::4} by the examples of photo-ionization cross-sections of all realized states with l=0. These figures show qualitative similarity of behavior of the cross-sections in the cases $r_{c}= 44.964 a.u.$ and  $ r_{c}= 55.052 a.u.$ and domination of the cross-sections with n=2. One can see a significant difference between the maximal values of the cross-sections with n=2 (about 2.70 a.u. and 0.75 a.u.) which correspond to these cases. This fact reflects the tendency of a significant decrease of the maximal values of the cross-sections for n=2 with an increase of the cut-off radius $r_{c}$.   Let us remind that Tab's 1 and 2 characterize the bound states of an electron in the potential $U_{c}(r)$ with the values of cut-off radius $r_{c}$ given above. The energies of the corresponding ground states approach the value of $-I_{H}$, where $ I_{H}=13.598 eV$ is the tabulated value of the isolated hydrogen-atom ionization potential (see for an example NIST Atomic Spectra Database), only when $r_{c}\\rightarrow \\infty$, i.e. when $N_e\\rightarrow 0$ or $T \\rightarrow \\infty$. Also, the energies of the ground states (for the electron densities and the temperatures observed) would be close to ($-I_{H}$) in the case where instead of $U_{c}(r)$ the potential $U(r; r_{c})$ would be used (see Fig.~1).  For each of the considered cut-off radii $r_{c}$ the existence is assumed here of a Boltzmann's distribution of the populations $N_{n,l}$ of the bound states, given in Tab 1 or 2, which exist in Eq. (\\ref{eq:kappaph}). Such a distribution is determined by the corresponding values of the total density of neutral hydrogen atoms $N_{a}$ and the temperature $T$. In accordance with \\cite{vit04} here it is taken that: $N_{a}=1.9\\cdot10^{19} cm^{-3}$ and $T=22980 K$ for $r_{c}= 44.964 a.u.$, and $N_{a}=3.4\\cdot10^{19} cm^{-3}$ and $T=17960 K$ for $ r_{c}= 55.052 a.u.$.  As one of the consequences, we have it that the total populations of groups of the states with same $n$ are equal to: $4.4\\cdot10^{17}$, $3.82 \\cdot10^{17}$ and $4.86 \\cdot10^{17} cm^{-3}$ for $n=2$, $3$ and $4$ in the first case, and $1.87 \\cdot10^{17}$,  $1.24 \\cdot10^{17}$, $1.44 \\cdot10^{17}$ and $1.85 \\cdot10^{17} cm^{-3}$ for $n=2$, $3$ , $4$ and $5$ in the second case.  \\subsection{The absorption coefficient: the results of the calculations}  In order to compare the obtained theoretical results with the experimental data from \\cite{vit04}, we have to take into account all the absorption processes which cannot be neglected in the considered hydrogen plasmas, i.e. the processes described by Eqs.~(\\ref{eq:ph}), (\\ref{eq:inbs}), (\\ref{eq:eH}) and (\\ref{eq:HH}). Therefore, when comparing our theoretical results with the experimental data from \\cite{vit04} we have to use the corresponding total spectral absorption coefficient $\\kappa_{tot}(\\lambda)$, namely \\begin{equation} \\label{eq:tot} \\kappa_{tot}(\\lambda) = \\kappa_{ph}(\\lambda) + \\kappa_{add}(\\lambda), \\end{equation} where $\\kappa_{ph}(\\lambda)\\equiv \\kappa_{ph}(\\lambda;N_{e},T)$ is given by Eq.~(\\ref{eq:kappaph}), and the member $\\kappa_{add}(\\lambda) \\equiv \\kappa_{add}(\\lambda;N_{e},N_{a},T)$ characterizes the contribution of absorption processes (\\ref{eq:inbs}), (\\ref{eq:eH}) and (\\ref{eq:HH}). Consequently, we have it that \\begin{equation} \\label{eq:add} \\kappa_{add}(\\lambda) = \\kappa_{ei}(\\lambda) + \\kappa_{ea}(\\lambda) + \\kappa_{ia}(\\lambda), \\end{equation} where $\\kappa_{ei}(\\lambda)\\equiv\\kappa_{ei}(\\lambda;N_{e},T)$, $\\kappa_{ea}(\\lambda)\\equiv\\kappa_{ea}(\\lambda; N_{e},N_{a},T)$ and $\\kappa_{ia}(\\lambda)\\equiv\\kappa_{ia}(\\lambda;N_{e},N_{a},T)$ are the partial spectral absorption coefficients, which are given above by Eqs.~(\\ref{eq:ei}), (\\ref{eq:ea}), (\\ref{eq:S}), (\\ref{eq:ia}) and (\\ref{eq:Kia}).  In accordance with the aims of this work the calculations of the total absorption coefficients  $\\kappa_{tot}(\\lambda)$ have been performed for both cases (short and long pulse) in a wide region of values of shifts ($\\Delta_{n}$) and broadenings ($\\delta_{n}$) of atomic levels with $n \\ge 2$. The calculations of $\\kappa_{tot}(\\lambda)$ cover the wavelength region $300nm \\le \\lambda \\le 500nm$. However, let us emphasize the fact that the values of the experimental absorption coefficient $\\kappa_{exp}(\\lambda)$ characterize not only the bound-free (photo-ionization) processes (\\ref{eq:ph}) and the said additional absorption processes (\\ref{eq:inbs}) - (\\ref{eq:HH}), but also the bound-bound (photo-excitation) processes, which are not considered in this work. Consequently, for the purpose of this work the region $\\lambda \\lesssim 450nm$ in the case of short pulse, and $\\lambda \\lesssim 425nm$ in the case of long pulse, (see Tabs.~1 and 2) has the real significance, where the considered photo-ionization processes dominate in comparison with photo-excitation ones. The results of calculations are shown in Figs.~5-10 together with the corresponding experimental values $\\kappa_{exp}(\\lambda)$ of the spectral absorption coefficient from \\cite{vit04}.  Figures 5, 6, 7 and 8 illustrate the results of the calculations of $\\kappa_{tot}(\\lambda)$ in the case when $\\Delta_{n} = const.$, while Figs.~9 and 10 present the calculations of $\\kappa_{tot}(\\lambda)$ in the case when $\\Delta_{n}$ decreases (relative to $\\Delta_{2}$) with increasing $n$.  The bottom and the top curves of $\\kappa_{tot}(\\lambda)$ in Figs.~5 and 6 illustrate strong influence of $\\Delta_{n}$ on the calculated total absorption coefficient: $\\Delta_{n}=\\delta_{n}=0.30eV$ and $\\Delta_{n}=\\delta_{n}=0.60eV$ for the short pulse; $\\Delta_{n}=\\delta_{n}=0.05eV$ and $\\Delta_{n}=\\delta_{n}=0.20eV$ for the long pulse. The groups of three curves, which lie between the corresponding bottom and top curves, demonstrate relatively small influence of $\\delta_{n}$ on the calculated values of $\\kappa_{tot}(\\lambda)$: $\\Delta_{n}=0.45eV$ and $\\delta_{n}=0.40eV$, $\\delta_{n}=0.45eV$ and $\\delta_{n}=0.50eV$ for the short pulse; $\\Delta_{n}=0.125eV$ and $\\delta_{n}=0.120eV$, $\\delta_{n}=0.125eV$ and $\\delta_{n}=0.130eV$ for the long pulse.  Also, the dashed curves on Figs.~5 and 6 demonstrate the behaviour of the spectral absorption coefficient $\\kappa_{add}(\\lambda)$, defined by Eqs. (\\ref{eq:add}) and (\\ref{eq:ei}) - (\\ref{eq:Kia}), in the case of short and long pulses respectively. One can see that the total contribution of electron-ion, electron-atom and ion-atom absorption processes, described by Eqs.~(\\ref{eq:inbs}), (\\ref{eq:eH}) and (\\ref{eq:HH}) indeed cannot be neglected in the considered cases.  Figures 7 and 8 show the curves of $\\kappa_{tot}(\\lambda)$ calculated with the values of $\\Delta_{n}$ and $\\delta_{n}$ which are treated as the optimal ones: $\\Delta_{n}= 0.455 eV$ and $\\delta_{n}= 0.625 eV$ for the short pulse, and $\\Delta_{n}=0.13 eV$ and $\\delta_{n}= 0.11 eV$ for the long one. In order to estimate the possible error due to such a choice of shifts and broadenings, the results of calculations are shown in Figs.~9 and 10 in the case when $\\Delta_{n}$ decreases with the increase of $n$, proportionally to the ionization energies of the corresponding atomic levels. Let us note the fact that calculated curves presented in these figures correspond to the optimal values of $\\Delta_{n=2}$ and $\\delta_{n=2}$. One can see that the calculated curves in Figs.~9 and 10 are very close to the calculated curves in Figs.~7 and 8, respectively. This fact is reflected in the values of $\\Delta_{n=2}$ and $\\delta_{n=2}$ which correspond to the curves in Figs.~9 and 10: $\\Delta_{n=2}=0.49 eV$ and $\\delta_{n=2}=0.65 eV$ for the short pulse, and $\\Delta_{n=2}=0.14 eV$ and $\\delta_{n=2}=0.12 eV$ for the long one.  Beside the curves $\\kappa_{tot}(\\lambda)$ and $\\kappa_{exp}(\\lambda)$, the curves $\\kappa_{p.th.}(\\lambda)$ are also presented in Figs.~7 and 8, which were obtained in \\cite{vit04} for the short and long pulse on the basis of the perturbation theory used in \\cite{gav01,vit04}. For the sake of correct interpretation of the presented data, let us note the fact that the values of $\\kappa_{p.th.}(\\lambda)$, similarly to $\\kappa_{exp}(\\lambda)$, characterize not only photo-ionization, but also the bound-bound (photo-excitation) processes, which are not considered in this work. One can see that in the case of strongly non-ideal plasmas (short pulse, $N_{e}=1.5\\cdot 10^{19} cm^{-3}$) the difference between $\\kappa_{exp}(\\lambda)$ and $\\kappa_{p.th.}(\\lambda)$ is so large that it justifies any effort towards development of an alternative method of calculation of the strongly non-ideal plasma absorption coefficient. However, in the case of long pulse ($N_{e}=0.65\\cdot 10^{19} cm^{-3}$) the considered plasma is located by its parameters in the lower part of the region of strong non-ideality, where the perturbation theory should give much better results. This fact is reflected in a significant reduction of the difference between $\\kappa_{exp}(\\lambda)$ and $\\kappa_{p.th.}(\\lambda)$.  Therefore it is important to check whether relation Eq.~(\\ref{eq:Deltan}) is valid also for $N_{e}$ close to $0.65\\cdot 10^{19} cm^{-3}$. Since in the case of constant shifts $\\Delta_{n}=0.455 eV$ and $0.130 eV$ for short and long pulses respectively, validity of Eq.~(\\ref{eq:Deltan}) means that $0.455/0.130=(1.5/0.65)^{3/2}$, which is satisfied with an accuracy better than $1\\%$. In the case of variable shift we have it that $\\Delta_{n=2}=0.49 eV$ and $0.12 eV$ for the short and long pulse respectively, and validity of Eq.~(\\ref{eq:Deltan}) means now that $0.49/0.14=(1.5/0.65)^{3/2}$, which is satisfied with the same accuracy.  This fact offers a possibility to determine $\\Delta_{n}$ or $\\Delta_{n=2}$ for any $N_{e}$ and $T$ from intervals $0.65 \\cdot 10^{19} cm^{-3} < N_{e} < 1.5\\cdot 10^{19} cm^{-3}$ and $1.8 \\cdot 10^{4} K \\le T \\le 2.3 \\cdot 10^{4} K$, and probably in some significantly wider regions, at least for $5\\cdot 10^{18} cm^{-3} \\lesssim N_e \\lesssim 5 \\cdot 10^{19}cm^{-3}$ and $1.6 \\cdot 10^{4} K \\lesssim T \\lesssim 2.5 \\cdot 10^{4} K$. Then, since it has been established that the influence of $\\delta_{n}$ is significantly weaker than the influence of $\\Delta_{n}$, we can determine $\\delta_{n}$ for any $N_{e}$ from those intervals, taking $\\delta_{n} = \\Delta_{n}$, as in the examples illustrated by the dashed curves on Figs.~\\ref{fig::7} and \\ref{fig::8}.  \\centerline{***}  On the grounds of all that has been said one can conclude that the presented method can already be used for calculations of the spectral absorption coefficients of dense hydrogen plasmas in the regions $N_{e}\\sim 10^{19} cm^{-3}$ and $T_{e}\\approx 2 \\cdot 10^{4}K$, as long as electron-$H^{+}$ and electron-$H(1s)$ inverse bremsstrahlung, negative ion $H^{-}(1s^{2})$ photo-detachment, $H(1s)$-$H^{+}$ absorption charge exchange and molecular ion $H_{2}^{+}(1\\Sigma^{+}_{g})$ photo-dissociation can be described as in this paper. Let us note the fact that with some modifications related to the atom photo-ionization processes, which should enable description of the influence of the atom core, the presented method can be applied to other kinds of laboratory dense hydrogen-like plasmas. Most of all, we mean alkali metals, helium and other rare-gas plasmas. Also, the obtained results can be of interest for some astrophysical plasmas, namely the plasma of the inner layers of solar atmosphere, as well as the plasmas of partially ionized layers of some other stellar atmospheres (for example some DA and DB white dwarfs).  Further development of the method described requires first of all an improvement of the procedure used, in order to replace the semi-empirical parameters with ones determined within the procedure itself. Another step would be including into consideration atom bound-bound (photo-excitation) processes which have been omitted here, and extending the region of the method's applicability to the long wavelengths."
    },
    "1208/1208.2625_arXiv.txt": {
        "abstract": "We constrain a stochastic background of primordial magnetic fields (PMF) by its contribution to the cosmic microwave background (CMB) anisotropy angular power spectrum with the combination of WMAP 7 year and South Pole Telescope (SPT) data. The contamination in the SPT data by unresolved point sources and by the Sunyaev Zeldovich (SZ) effect due to galaxy clusters has been taken into account as modelled by the SPT collaboration. With this combination of WMAP 7 yr and SPT data, we constrain the amplitude Gaussian smoothed over 1 Mpc scale of a stochastic background of non-helical PMF to $B_{\\rm 1 Mpc}<3.5$ nG at 95\\% confidence level, improving on previous bounds. Our analysis shows that SPT data up to $\\ell=3000$ bring an improvement of almost a factor two with respect to results with previous CMB high-$\\ell$ data. We then discuss the forecasted impact from unresolved points sources and SZ effect for {\\sc Planck} capabilities in constraining PMF. ",
        "introduction": "Current CMB anisotropy measurements lead to upper limits on the amplitude of a stochastic background of primordial magnetic fields generated before nucleosynthesis \\cite{yamazaki,yamazakilast,PF,Yamazaki2012,shawlewis2}. Indeed, a stochastic background of PMF generates all types of magnetized linear perturbations \\cite{PFP,SL}: scalar \\cite{KL,yamazaki,KR,GK,FPP,PFP,SL,BoCa}, vector \\cite{SB,MKK,lewis} and tensor \\cite{DFK,MKK,CDK} and all these contribute to the CMB anisotropy pattern in temperature and polarization. CMB constraints on PMF with the angular power spectrum agree with those from their effect on the reionization epoch \\cite{MiniatiConstraints}. PMF modelled as a fully inhomogeneous component have also a fully non-Gaussian contribution to CMB anisotropies with a non zero higher statistical moments, which can be used as useful probes, such as the magnetized bispectrum \\cite{CFPR,SS} and the magnetized trispectrum \\cite{Trivedi}.  In our previous works \\cite{FPP,PFP,PF} we have refined the computation of magnetized CMB anisotropies. In Ref. \\cite{PF} we have computed the constraints coming from CMB data by WMAP7 in combination with data from ACBAR \\cite{ACBAR}, QUaD \\cite{QUAD} and BICEP \\cite{BICEP} updating previous investigations \\cite{yamazaki,yamazakilast,shawlewis2}.  In this work we use the publicly available CMB anisotropy data at high multipoles as those from the South Pole Telescope (SPT) \\cite{SPT2011,Reichardt} to further constrain a stochastic background of PMF. Constraints on PMF from CMB anisotropies at high multipoles, $\\ell \\sim 3000$, are not a straightforward extension of those derived at larger angular scales. Small angular scale are in fact polluted by extragalactic contamination \\cite{Reichardt,Millea,PaolettiADFLDP} and secondary anisotropies, such as Sunyaev-Zeldovich \\cite{TSZOriginal,KSZOriginal}. In order to fully exploit small scale CMB data to constrain PMF it is necessary to model the residual foreground contamination to the angular power spectrum. ",
        "conclusions": "We have derived the constraints on a stochastic background of PMF by using the CMB temperature anisotropy measurements at high multipoles by SPT. This study is motivated by the fact that  the PMF contribution to CMB anisotropies is not suppressed by Silk damping as the primary anisotropies.  In order to not introduce biases in the magnetic parameter constraints we need to consider the contamination by astrophysical residuals of the SPT data. The dominant contributions are given by unresolved point sources and in particular radio and infrared galaxies and by galaxy clusters. We have considered both Poissonian and clustering terms for point sources and the SZ effect for the galaxy cluster contribution. To model the contributions to the angular power spectrum of the three signals we have used the templates provided by the SPT collaboration \\cite{SPT2011,SPTsite}. We performed a MCMC analysis with the eleven cosmological, magnetic and astrophysical parameters and we constrain $B_{1 {\\rm Mpc}} < 3.5$ nG. The results do not show any strong degeneracy between magnetic and astrophysical parameters which is compatible with the multipole range of SPT data ($\\ell_{max}\\sim 3000$) used. Comparing these results with the previous constraints with data by WMAP7, ACBAR, QUAD and BICEP \\cite{PF,shawlewis2}, which were of the order of $B_{1 {\\rm Mpc}} < 5$ nG, we note a drastic improvement in the constraint on $B_{1 {\\rm Mpc}}$ with the use of SPT data. We have shown how the current CMB constraints for our choice of $k_D$ are by far tighter than those derived from BBN for all the $n_B$ considered here.  We have also updated the expected constraints from {\\it Planck} by including the astrophysical contamination at small angular scales following the treatment in \\cite{PaolettiADFLDP}. The results we obtained show a (expected) degradation of the constraints on PMF due to the presence of extragalactic contributions: $B_{1 {\\rm Mpc}} < 3.6$ nG, compared to the previous constrain: $B_{1 {\\rm Mpc}} < 2.4$ nG (obtained taking into account only noise and sensitivity). The results presented here confirmed a previously noted trend  which prefer negative $n_B$. Since $n_B>0$ is mainly related to causal generation mechanism, we have shown again how causal fields are allowed with an amplitude much smaller than the nGauss level.  \\vspace{1cm}  {\\bf Acknowledgements} We acknowledge support by PRIN MIUR 2009 grant n. 2009XZ54H2 and ASI through ASI/INAF Agreement I/072/09/0 for the Planck LFI Activity of Phase E2."
    },
    "1208/1208.3309_arXiv.txt": {
        "abstract": "The availability of photometric imaging of several thousand galaxies with the {\\it Spitzer Space Telescope} enables a mid-infrared calibration of the correlation between luminosity and rotation in spiral galaxies.  The most important advantage of the new calibration in the $3.6\\mu$m band, IRAC ch.1, is photometric consistency across the entire sky.  Additional advantages are minimal obscuration,  observations of flux dominated by old stars, and sensitivity to low surface brightness levels due to favorable backgrounds. Through Spitzer cycle 7 roughly 3000 galaxies had been observed and images of these are available at the Spitzer archive. In cycle 8 a program called {\\it Cosmic Flows with Spitzer} has been initiated that will increase by 1274 the available sample of spiral galaxies with inclinations greater than 45 degrees from face-on suitable for distance measurements.  This paper describes procedures based on the photometry package Archangel that are being employed to analyze both the archival and the new data in a uniform way. We give results for 235 galaxies, our calibrator sample for the Tully-Fisher relation.  Galaxy magnitudes are determined with uncertainties held below 0.05 mag for normal spiral systems.  A subsequent paper will describe the calibration of the [3.6] luminosity$-$rotation relation. ",
        "introduction": "The Cosmic Microwave Background (CMB) temperature dipole \\citep{1996ApJ...473..576F} is usually interpreted as a motion of our Galaxy of over 600~\\kms\\ but the considerable majority of the posited motion is developed on large scales with origins that are still poorly understood.  Our overarching goal is to measure distances, hence parse departures from the mean Hubble expansion, on scales extending to 200~Mpc.  We are gathering distance measures from a multitude of methods and contributors within a program that we are calling {\\it Cosmic Flows} (see Appendix).  Of particular importance for us are distances accrued from the correlation between the rotation rate of a galaxy and its luminosity, the Tully-Fisher Relation: TFR \\citep{1977A&A....54..661T}.  There are methodologies that provide distance estimates that are individually more accurate but an abiding advantage of the TFR is applicability to a large fraction of all galaxies over a wide range of environments and distances.  There is the prospect of utilizing the TFR to obtain distances to several tens of thousands of galaxies out to redshift $z\\sim0.05$ ($\\sim$ 200 Mpc) .\\\\  As steps toward the accumulation of an appropriate data set of distances, our Cosmic Flows Large Program on the 100m Green Bank Telescope and complementary southern  observations on the Parkes Telescope \\citep{2009AJ....138.1938C,2011MNRAS.414.2005C} are providing us with quality rotation information from HI line profiles, and the merging of our own and literature optical photometry \\citep{2011MNRAS.415.1935C} provides the other element that has permitted a modern re-calibration of the TFR at an optical band \\citep{2012ApJ...749...78T}.  The data accumulated in support of this program, including the new material to be discussed in this paper, are made available at EDD, the Extragalactic Distance Database, accessed online at http://edd.ifa.hawaii.edu \\citep{2009AJ....138..323T}.\\\\  It has long been appreciated that photometry in the infrared may offer advantages because of reduced extinction and because infrared flux arises in large measure from old stars that should dominate the inventory of baryonic mass \\citep{1979ApJ...229....1A}.  Photometry at $K_s$ band from the Two Micron All Sky Survey (2MASS) has been used with the TFR \\citep{2002A&A...396..431K}.  However a major concern with observations in the infrared from the ground is high and variable sky foreground.  Much of the flux from galaxies lies in extended components with surface brightnesses that are well below the ground based sky level. Flux at the extremities of galaxies is lost and very low surface brightness galaxies are not even seen. \\\\  Observations from space removes the problem of the high contamination by the Earth's atmosphere.  We have initiated a sub-program that we call {\\it Cosmic Flows with Spitzer} (CFS) with NASA's {\\it Spitzer Space Telescope} \\citep{2004ApJS..154....1W}.  Observations have begun in cycle 8 during the post-cryogenic period to obtain wide-field images of galaxies with IRAC, the InfraRed Array Camera, in Channel 1 (ch.1).  The present paper describes our reduction and photometry analysis procedures. We will discuss the transformation steps from raw Spitzer Post Basic Calibrated Data (PBCD) obtained from the Spitzer archive to the parameter needed for the TFR calibration: apparent [3.6] magnitudes.  The photometry is carried out with a Spitzer-adapted version of Archangel \\citep{2007astro.ph..3646S,2012PASA...29..174S}. We will discuss the corrections to be made to apparent magnitudes and conclude with a discussion of uncertainties.  A subsequent paper will discuss the calibration of the TFR at $3.6 \\mu$m. ",
        "conclusions": "Our photometric procedures for the semi-automated analysis of Spitzer IRAC channel 1 data at $3.6 \\mu$m have been described.  The galaxy surface photometry is carried out with the Archangel software \\citep{2007astro.ph..3646S,2012PASA...29..174S} adapted for Spitzer data input.  Already material is available for some 3000 galaxies from the Spitzer Heritage Archive and our {\\it Cosmic Flows with Spitzer} program will supply information for an additional 1274 galaxies. The 235 galaxies analyzed in the course of this paper will be used in a subsequent paper for the calibration of the mid-infrared Tully-Fisher relation.\\\\  The final goal of our project is to measure distances, hence map peculiar velocities, across the local universe within $10,000$~km s$^{-1}$ using the correlation between galaxy luminosities and their rotation rates.  We have demonstrated the ability to use {\\it Spitzer Space Telescope} mid-IR data to perform surface photometry with a relatively high accuracy.   No correlation is found between magnitude uncertainties and other important galaxy parameters such as inclination, apparent area, or semi-major axis.  We conclude that, after all corrections, uncertainties on magnitudes are of the order $\\pm 0.05$ for the regular spiral galaxies at the heart of our project.  These uncertainties are small compared with the overall scatter in the TFR. Low surface brightness galaxies or very irregular ones require special attention but these classes of galaxies are not of principal interest to us. \\\\  For our purposes, the advantages of mid-infrared photometry from space include minimization of both galactic and internal obscuration issues, very low backgrounds, and source fluxes dominated by old stellar populations that are good representatives of the baryonic mass.  The most outstanding advantage, though, is the integrity and consistency of the photometry in all quadrants of the sky.  The Extragalactic Distance Database, EDD, contains HI profile information that provides useful line widths for over 11,000 galaxies.  Ongoing Spitzer observations are providing the complementary photometric information required for a dense, detailed map of structure and motions in the near part of the Universe.\\\\  \\bigskip  Acknowledgements. Comparisons between alternative analysis procedures have been facilitated by data made available by Tom Jarrett, Mark Seibert representing CHP, the {\\it Carnegie Hubble Program}, and Kartik Sheth, representing S4G, the {\\it Spitzer Survey of Stellar Structure in Galaxies}.  Thanks to James Shombert for the development and support of the Archangel photometry package.   The referee Michael Pohlen encouraged us to write the overview of {\\it Cosmic Flows} which constitutes the Appendix and was otherwise very helpful. NASA through the Spitzer Science Center provides support for CFS, {\\it Cosmic Flows with Spitzer},  cycle 8 program 80072.  In addition to the authors, CFS co-investigators are Wendy Freedman, Tom Jarrett, Barry Madore, Eric Persson, Mark Seibert, and Ed Shaya. RBT receives support for aspects of this program from the US National Science Foundation with award AST-0908846.  \\bigskip\\noindent{\\bf Appendix: Cosmic Flows Program Overview}\\\\  {\\it Cosmic Flows} may have as many arms as an octopus.  At its core is a collaboration between Courtois and Tully to obtain accurate distances to galaxies.   A major part of the program involves exploitation of the TFR.   Activities in this regard began with the accumulation of HI profiles for the necessary kinematic information within the {\\it Cosmic Flows Large Program} using the US National Radio Astronomy Green Bank Telescope \\citep{2009AJ....138.1938C, 2011MNRAS.414.2005C} and the accumulation of optical photometry for the necessary magnitude and inclination information using the University of Hawaii 2.24m telescope \\citep{2011MNRAS.415.1935C}.   The present extension embarks on complementing the optical photometry with mid-infrared photometry.  The current paper describes analysis procedures developed within the core program but {\\it Cosmic Flows with Spitzer} embraces the larger team identified in the acknowledgement.  The next paper in this series will include most of the CFS team in a discussion of the TFR calibration with Spitzer photometry.\\\\  Near, intermediate, and far TFR samples in the {\\it Cosmic Flows} program were described by \\citet{2011MNRAS.414.2005C}.  The `near' sample is intended to achieve dense coverage of a volume extending to 3300~\\kms\\ with inclusion of all galaxies typed later than Sa that are brighter than $M_K =-21$, inclined greater than $45^{\\circ}$, and not obscured, disrupted, or confused.  The `intermediate' sample is drawn from flux and color limits applied to an Infrared Astronomical Satellite redshift survey \\citep{2000MNRAS.317...55S}.  The flux limit at $60~\\mu$m is 0.6 Jy, the color criterion to separate normal spirals from active nuclei is a ratio of $100~\\mu$m to $60~\\mu$m flux greater than one, there is a velocity cutoff at 6000~\\kms, and there is the same inclination restriction as with the near sample.  By contrast, the `far' sample is restricted to extreme edge-on systems drawn from Flat Galaxy catalogues \\citep{1999BSAO...47....5K, 2004BSAO...57....5M}.  Candidates in the sample that lie at declinations accessible to Arecibo Telescope have velocities extending to 15,000~\\kms.  These are our well defined samples.  In addition we derive distances to all other suitably observed galaxies.  Generally the information for the additional systems comes from archives.  In all, presently good data are available for about 7500 appropriate galaxies.\\\\  A quite separate and active component of {\\it Cosmic Flows} is a program with {\\it Hubble Space Telescope} to obtain Tip of the Red Giant Branch distances to nearby, spatially resolved galaxies \\citep{2006AJ....132.2729M, 2007ApJ...661..815R, 2009AJ....138..332J}. Exquisite distances (5\\% accuracy) are available for approaching 300 galaxies within $\\sim 10$~Mpc.\\\\  Distances for {\\it Cosmic Flows} encompass measures by other methodologies discussed in the literature.  Foremost among these are Cepheid Period-Luminosity Relation, Surface Brightness Fluctuation, Fundamental Plane, and Supernova Ia procedures.  The diverse material is drawn together in EDD, the Extragalactic Distance Database\\footnote{http://edd.ifa.hawaii.edu} \\citep{2009AJ....138..323T}.  EDD goes beyond the compilation of catalogs relevant to extragalactic distances to include redshift catalogs, that with various levels of completion describe the distribution of galaxies in the local universe, and group catalogs, that help identify entities where averaging over velocities or distances is reasonable.  The first assembly of distances in this program \\citep{2008ApJ...676..184T} has now been given the name {\\it Cosmicflows-1}.  A core team is now involved in the assembly of {\\it Cosmicflows-2} \\citep{2012ApJ...749...78T, 2012ApJ...749..174C}.\\\\  The holy grail of {\\it Cosmic Flows} is the use of distances to determine peculiar velocities and, subsequently, mass fluctuations.  Peculiar velocities are departures from the cosmic mean expansion and it is assumed that they arise due to density irregularities.  Two regimes require separate attention.  The high density environments in and around collapsed halos are at the extreme of non-linear dynamics.  Within the collaboration we have developed Numerical Action Methods that provide an optimal description of the distribution of mass affecting galaxies on curved orbits on first approach to an attractor \\citep{1995ApJ...454...15S, 2001ApJ...554..104P, 2011arXiv1105.5596P}.  The other extreme is the regime of linear dynamics.  A procedure we have used that is appropriate with redshift data sets of $10^5$ or more objects is based on the action principle \\citep{2010ApJ...709..483L}.  However the methodology that most interests us starts with Wiener filtering of the peculiar velocity field resulting in descriptions of the density field independent of information provided by redshift surveys \\citep{1995ApJ...449..446Z, 2012ApJ...744...43C}. The current density field can be mapped back to initial conditions that are then the starting point for constrained simulations that attempt to approximate the observed universe with a computer model \\citep{2003ApJ...596...19K, 2010arXiv1005.2687G, 2012arXiv1205.4627C}.  \\clearpage"
    },
    "1208/1208.4529_arXiv.txt": {
        "abstract": "Magnetic field generation on scales large compared with the scale of the turbulent eddies is known to be possible via the so-called $\\alpha$ effect when the turbulence is helical and if the domain is large enough for the $\\alpha$ effect to dominate over turbulent diffusion. Using three-dimensional turbulence simulations, we show that the energy of the resulting mean magnetic field of the saturated state increases linearly with the product of normalized helicity and the ratio of domain scale to eddy scale, provided this product exceeds a critical value of around unity. This implies that large-scale dynamo action commences when the normalized helicity is larger than the inverse scale ratio. \\blue{ Our results show that the emergence of small-scale dynamo action does not have any noticeable effect on the large-scale dynamo. } Recent findings by Pietarila Graham et al.\\ (2012, Phys.\\ Rev.\\ E85, 066406) of a smaller minimal helicity may be an artifact due to the onset of small-scale dynamo action at large magnetic Reynolds numbers. However, the onset of large-scale dynamo action is difficult to establish when the kinetic helicity is small. Instead of random forcing, they used an ABC-flow with time-dependent phases. We show that such dynamos saturate prematurely in a way that is reminiscent of inhomogeneous dynamos with internal magnetic helicity fluxes. Furthermore, even for very low fractional helicities, such dynamos display large-scale fields that change direction, which is uncharacteristic of turbulent dynamos. ",
        "introduction": "The origin of magnetic fields in astrophysical bodies like the Earth, the Sun and galaxies is studied in the field of dynamo theory. The temporal variation and strength of those fields rules out a primordial origin, through which the magnetic field would have been created in the early Universe. For magnetic fields with energies of the equipartition value, \\ie the turbulent kinetic energy of the medium, the primordial hypothesis explains their strength after creation, but falls short of explaining how the field outlives billions of years of resistive decay \\cite{Parker1979book}.  In dynamo theory, astrophysical plasmas are considered sufficiently well conducting fluids where the inertia of the charge-carrying particles can be neglected. In this approximation the equations of magnetohydrodynamics (MHD) provide an adequate model of the medium. In this framework it has been studied under which conditions magnetic fields of equipartition strength and scales larger than the turbulent motions are created and sustained \\cite{MoffattBook1978}.  A successful theoretical model describing the dynamo's behavior is the mean-field theory. It relates the small-scale turbulent motions to the mean magnetic field via the so-called $\\alpha$ effect, which provides the energy input via helical turbulent forcing. During the kinematic phase, \\ie negligible back reaction of the magnetic field on the fluid, the $\\alpha$ effect gives a positive feedback on the large-scale magnetic field, which results in its exponential growth. Already the consideration of the kinematic MHD equations with negligible Lorentz force sheds light on the growth rate of the different modes of the magnetic field during the kinematic phase. In the kinematic phase the growth rate $\\lambda$ at wave number $k$ is given by \\cite{MoffattBook1978} \\EQ \\label{eq: lambda kinetic} \\lambda  = \\alpha k - \\etaT k^{2} = (\\Ca-1) \\etaT k^{2}, \\EN where $\\Ca=\\alpha/(\\etaT k)$ is the relevant dynamo number for the $\\alpha^2$ dynamo, $\\alpha$ is the $\\alpha$ coefficient which is proportional to the small-scale kinetic helicity, and $\\etaT = \\eta + \\eta_{\\rm t}$ is the sum of molecular and turbulent magnetic diffusivity. Clearly, dynamo action occurs when $|\\Ca|>\\Cacrit$, where the onset condition is $\\Cacrit=1$. Standard estimates for isotropic turbulence in the high conductivity limit \\cite{Radler1980book,MoffattBook1978} yield $\\alpha\\approx-(\\tau/3)\\bra{\\oo\\cdot\\uu}$ and $\\etat\\approx(\\tau/3)\\bra{\\uu^2}$, where $\\tau$ is the correlation time of the turbulence, $\\oo=\\nab\\times\\uu$ is the vorticity and $\\uu$ is the velocity in the small-scale fields. Here, $\\langle . \\rangle$ denotes a volume average. Using $\\etat\\gg\\eta$, we have \\EQ \\Ca\\approx-\\bra{\\oo\\cdot\\uu}/(k\\bra{\\uu^2}). \\EN It is convenient to define $\\bra{\\oo\\cdot\\uu}/(\\kf\\bra{\\uu^2})$ as the normalized kinetic helicity, $\\epsf$, so $\\Ca\\approx-\\epsf\\kf/k$. This scaling implies that the critical value of the normalized helicity $\\epsf$ scales inversely proportional to the scale separation ratio, i.e.\\ $\\epsf^{\\rm crit} \\propto(\\kf/k)^{-1}$, where $k\\ll\\kf$ is the wave number of the resulting large-scale magnetic field. This wave number can be equal to $k=k_1\\equiv2\\pi/L$, which is the smallest wave number in a periodic domain of size $L$.  In summary, the critical dynamo number $\\Cacrit$, which decides between growing or decaying solutions of the large-scale dynamo (LSD), is proportional to the product of normalized helicity $\\epsf$ and scale separation ratio $\\kf/k$. Therefore, the amount of helicity needed for the LSD is inversely proportional to the scale separation ratio, and not some higher power of it. It should be noted that the {\\em normalized} kinetic helicity $\\epsf$ used here is not the same as the {\\em relative} kinetic helicity, $\\tilde\\epsf=\\bra{\\oo\\cdot\\uu}/(\\orms\\urms)$. The two are related to each other via the relation \\EQ \\tilde\\epsf/\\epsf=(\\kom/\\kf)^{-1}, \\EN where $\\kom\\approx\\orms/\\urms$ is inversely proportional to the Taylor microscale. Here, the subscripts rms refer to root-mean-square values. For small Reynolds numbers, $\\kom$ provides a useful estimate of the wave number $\\kf$ of the energy-carrying eddies. In contrast, for large Reynolds numbers $\\Rey$, we expect $\\kom/\\kf$ to be proportional to $\\Rey^{1/2}$, so $\\tilde\\epsf$ decreases correspondingly while $\\epsf$ remains unchanged.  To understand the saturation of a helical dynamo, it is important to understand the relation between the resulting large-scale field and the associated small-scale field. Indeed, the growth of the large-scale field is always accompanied by a growth of small-scale magnetic field. Small-scale here means the scale of the underlying turbulent motions, which drive the dynamo. Conservation of total magnetic helicity causes a build-up of magnetic helicity at large scales and of opposite sign at small scales \\cite{Seehafer1996,Ji1999}. As the dynamo saturates, the largest scales of the magnetic field become even larger, which finally leads to a field of a scale that is similar to that of the system itself. This can be understood as being the result of an inverse cascade, which was first predicted based on closure calculations \\cite{Frisch-Pouquet-Leorat-1975-JFluidMech}.  If the domain is closed or periodic, the build-up of small-scale magnetic helicity causes the $\\alpha$ effect to diminish, which marks the end of the exponential growth and could occur well before final saturation is reached. The dynamo then is said to be catastrophically quenched and, in a closed or periodic system, the subsequent growth to the final state happens not on a dynamical timescale, but on a resistive one. Quenching becomes stronger as the magnetic Reynolds number increases, which, for astrophysically relevant problems, means a total loss of the LSD within the timescales of interest. In the case of open boundaries magnetic helicity fluxes can occur, which can alleviate the quenching and allow for fast saturation of the large-scale magnetic field \\cite{BlackmanField2000a,BlackmanField2000b,Kleeorin_etal2000,ssHelLoss09}.  In a recent publication \\cite{GrahamBlackmanMinuteHelicity12} it was argued that for periodic boundaries the critical value of $\\epsf$ for LSD action to occur decreases with the scale separation ratio like $\\epsf^{\\rm crit} \\propto (\\kf/k_1)^{-3}$. Their finding, however, is at variance with the predictions made using equation \\eqref{eq: lambda kinetic}, which would rather suggest a dependence of $\\epsf^{\\rm crit} \\propto (\\kf/k_1)^{-1}$ with $\\Cacrit = 1$. This discrepancy could be a consequence of the criterion used in \\cite{GrahamBlackmanMinuteHelicity12} for determining $\\Cacrit$. The authors looked at the growth rate of the magnetic field after the end of the kinematic growth phase, but only at a small fraction of the resistive time. Therefore their results might well be contaminated by magnetic fields resulting from the small-scale dynamo (SSD). Earlier simulations \\cite{Brandenburg2009ApJ} have demonstrated that for $\\Rm\\ge100$, the growth rate of the helical LSD approaches the well-known scaling of the nonhelical SSD with $\\lambda\\propto\\Rey^{1/2}$, which corresponds to the turnover rate of the smallest turbulent eddies \\cite{Schekochihin_etal2004ApJ,Haugen_etal2004PhRvE}.  Given that the LSD is best seen in the nonlinear regime \\cite{BrandInverseCascade2001}, we decided to determine $\\Cacrit$ from a bifurcation diagram by extrapolating to zero. In a bifurcation diagram, we plot the energy of the mean or large-scale field versus $\\Ca$. Simple considerations using the magnetic helicity equation applied to a homogeneous system in the steady state show that the current helicity must vanish \\cite{BrandInverseCascade2001}. In a helically driven system, this implies that the current helicity of the large-scale field must then be equal to minus the current helicity of the small-scale field. For a helical magnetic field, the normalized mean square magnetic field, $\\bra{\\meanBB^2}/\\Beq^2$, is approximately equal to $\\Ca-\\Cacrit$. Here, $\\Beq = (\\mu_{0}\\mean{\\rho})^{1/2}\\urms$ is the equipartition value of the magnetic field, $\\mu_{0}$ is the vacuum permeability, and $\\mean{\\rho}$ is the mean density. Again, since $\\Cacrit\\approx1$ and $\\Ca\\approx\\epsf\\kf/k_1$, this suggests that the LSD is excited for $\\epsf>(\\kf/k_1)^{-1}$ rather than some higher power of $\\kf/k_1$. This is a basic prediction that has been obtained from nonlinear mean-field dynamo models that incorporate magnetic helicity evolution \\cite{BlackmanBrandenb2002ApJ} as well as from direct numerical simulations in the presence of shear \\cite{KapylaBrandenb2009ApJ}. It is important to emphasize that mean field dynamo theory has been criticized on the grounds that no $\\alpha$ effect may exist in the highly nonlinear regime at large magnetic Reynolds numbers \\cite{CH09}. This is however in conflict with results of numerical simulations using the test-field method \\cite{BRRS08} showing that $\\alpha$ effect and turbulent diffusivity are both large, and that only the difference between both effects is resistively small. Another possibility is that the usual helical dynamo of $\\alpha^2$ type may not be the fastest growing one \\citep{BCR05}. This is related to the fact that, within the framework of the Kazantsev model \\cite{Kaz68} with helicity, there are new solutions with long-range correlations \\citep{S99,BS00}, which could dominate the growth of a large scale field at early times. The purpose of the present paper is therefore to reinvestigate the behavior of solutions in the nonlinear regime over a broader parameter range in the light of recent conflicting findings \\cite{GrahamBlackmanMinuteHelicity12}. ",
        "conclusions": "In this paper we have studied the simplest possible LSD and have investigated the dependence of its saturation amplitude on the amount of kinetic helicity in the system. We recall that the case of a periodic domain has already been investigated in some detail \\cite{FieldBlackman2002ApJ,Subramanian2002}, and that theoretical predictions in the case with shear \\cite{BlackmanBrandenb2002ApJ} have been verified numerically for fractional helicities \\cite{KapylaBrandenb2009ApJ}. Yet the issue has now attracted new interest in view of recent results suggesting that, in the limit of large scale separation, the amount of kinetic helicity needed to drive the LSD might actually be much smaller than what earlier calculations have suggested \\cite{GrahamBlackmanMinuteHelicity12}. This was surprising given the earlier confirmations of the theory. As explained above, the reason for the conflicting earlier results may be the fact that the LSD cannot be safely isolated in the linear regime, because it will be dominated by the SSD or, in the case of the ABC-flow dynamo, by some other kind of dynamo that is not due to the $\\alpha$ effect. Furthermore, as already alluded to in the introduction, there can be solutions with long-range correlations that could mimic those that are not due to the $\\alpha$ effect. Within the framework of the Kazantsev model \\cite{Kaz68}, the solutions to the resulting Schr\\\"odinger-type equation can be described as bound states. The addition of kinetic helicity leads to new solutions with long-range correlations as a result of tunneling from the SSD solutions \\citep{S99,BS00,BCR05}. Indeed, it has been clear for some time that large-scale magnetic fields of the type of an $\\alpha^2$ dynamo become only apparent in the late saturation of the dynamo \\cite{BrandInverseCascade2001}. This is especially true for the case of large values of $\\Rm$ when the mean field develops its full strength while the rms value of the \\blue{ small-scale field due to SSD action remains approximately unchanged as $\\Rm$ increases; see \\Fig{pncomp_flucts}. }  While there will always remain some uncertainty regarding the application to the much more extreme astrophysical parameter regime, we can now rule out the possibility of surprising effects within certain limits of $\\Rm$ and $\\Rey$ below 740, and scale separation ratios below 80. In stars and galaxies, the scale separation ratio is difficult to estimate, but it is hardly above the largest value considered here. This ratio is largest in the top layers of the solar convection zone where the correlation length of the turbulence is short ($1\\Mm$) compared with the spatial extent of the system ($100\\Mm$).  Of course, the magnetic Reynolds numbers in the Sun and in galaxies are much larger than what will ever be possible to simulate. Nevertheless, the results presented here show very little dependence of the critical value of $\\Ca$ on $\\Rm$. For $\\Pm=1$, for example, we find $\\Cacrit=1.2$ for $\\Rm\\approx6$ and $\\Cacrit=1.5$ for $\\Rm\\approx600$. On the other hand, for larger values of $\\Pm$, the value of $\\Cacrit$ can drop below unity ($\\Cacrit=0.9$ for $\\Pm=100$). While these changes of $\\Cacrit$ are theoretically not well understood, it seems clear that they are small and do not provide support for an entirely different scaling law, as anticipated in recent work \\cite{GrahamBlackmanMinuteHelicity12}."
    },
    "1208/1208.4003_arXiv.txt": {
        "abstract": "{  We present results of~nearly six years of~spectroscopic observations of~the~B[e] star V2028~Cyg. The~presence of~the~cold-type absorption lines combined with~a~hot-type spectrum indicate the~binarity of~this object. Since B[e] stars are embedded in~an~extended envelope, the~usage of~common stellar atmosphere models for~the~analysis is quite inappropriate. Therefore, we focus on~the~analysis of~the~long-term spectral line variations in~order to~determine the~nature of~this object. We present the~time dependences of~the~equivalent width and~radial velocities of~the~H$\\alpha$ line, [\\ion{O}{i}] 6300~\\AA, \\ion{Fe}{ii} 6427, 6433, and~6456~\\AA \\, lines. The~bisector variations and~line intensities are shown for~the~H$\\alpha$ line. The~radial velocities are also measured for~the~absorption lines of~the~K component. No~periodic variation is found. The~observed data show correlations between the~measured quantities, which can be used in~future modelling. } ",
        "introduction": "V2028~Cyg (MWC~623) belongs to~a~rather heterogeneous group called B[e] stars. This group includes such dissimilar objects as compact planetary nebulae, Herbig stars, and supergiants or~symbiotic stars \\citep{Lamers}. A~large number of~B[e] stars still belong to~this \"unclassified\" subgroup, which is also the~case for~V2028~Cyg. It is difficult to~assign this object to~a~specific subgroup because of~the~presence of~a~circumstellar matter. \\cite{miros07} found that many unclassified objects display common spectral signatures. He denoted this group as FS~CMa stars, since this star can be taken as a~good prototype for~these objects. He included in~this group V2028~Cyg. Current observation analysis shows that~they are nonsupergiants. FS~CMa objects are probably more evolved stars, but~are either close to the~main sequence, or~still on~it \\citep{miros00, miros07}. \\cite{miros07} suggested that~a~binarity could explain the~observed properties of~these stars.  The commonly used model of~B[e] stars was proposed by~\\cite{zickgraf85}. The~B[e] phenomenon arises in~a~two-component stellar wind with~a~dusty ring. A~slow, cool, high density wind occurs around the~equatorial plane. This wind forms an~outflowing disc, where the~emissions of~both the~Balmer series and~low-ionized metals originate. The~polar areas are~occupied by~a~CAK \\citep{CAK} wind of~a~low density and~high velocity. The~gas in~this outflow is hot and~the~emission lines come from~metals in~higher ionization stages.  V2028~Cyg was first noted as~an~emission-line star by~\\cite{merrill42}. \\cite{allsw76} discovered its photometrical variability (amplitude 0.5~mag in~filter V). \\cite{bergner95} then presented the~amplitude $0.2-0.4$~mag. \\cite{allen73} analysed photometric measurements from~the~filters H, K, and~L and~reported an~infrared excess that~he explained in~terms of~a~dust envelope. \\cite{allen74} then noted the~connection between~the~excess and~forbidden emission lines. A~few years later, \\cite{allsw76} published a~spectral analysis of~the~group of~peculiar Be stars with~infrared excesses, including V2028~Cyg. They identified permitted and~forbidden emission lines, such as~Balmer lines, \\ion{Fe}{ii}, [\\ion{O}{i}], and~[\\ion{Fe}{ii}], which~dominate the~emission spectrum.  V2028~Cyg was noted for~the~first time as~a~binary by~\\cite{Arkhipova82}. They suggested spectral the~types B8 and~M1III for~the~system from the~analysis of~spectrophotometric observations. However, they used assumptions that~were later found to~be~incorrect. High resolution spectra enabled \\cite{Zickgraf89} to~identify many absorption lines of mainly neutral metals, such as \\ion{Fe}{i}, \\ion{Ti}{i}, \\ion{V}{i}, \\ion{Ca}{i}, and~\\ion{Li}{i} (also detected later by~\\citealt{corporon} and~\\citealt{Zickgraf01}). By~fitting theoretical spectral energy distribution to~photometric data, they suggested that~V2028~Cyg is a~binary composed of a~B2 dwarf and~a~K2 giant. From~comparison with~spectra of~normal stars, \\cite{Zickgraf01} refined the~component classification to~B4III and~K2Ib-II. Multicolour optical and~near-infrared photometric observations led \\cite{bergner95} to~derive the~type K7III for~the~cool component. The~primary component falls on~the~HR diagram into the region where other FS~CMa stars are also found \\citep{miros07}. Its position close to~the~region occupied by~classical Be stars, provides a~good opportunity to study the~connection between the~Be and~B[e] phenomenon.  \\cite{Zickgraf01} derived a~kinematical distance towards V2028~Cyg using the~average radial velocity ($2.0^{+0.6}_{-0.3}$~kpc) and~his luminosity estimate for~both components ($2.4^{+1.4}_{-0.9}$~kpc).  The~optical emission spectrum consists of~a~large number of~singly ionized metallic lines (\\ion{Fe}{ii}, \\ion{Ti}{ii}, \\ion{Cr}{ii}, \\ion{Si}{ii}) and~is dominated by~the~hydrogen lines of~the~Balmer series. V2028~Cyg possesses an~extended gaseous and~dusty envelope, which is indicated by~several spectral features (IR excess, forbidden lines, intensive Balmer emission). \\cite{marston} unsuccessfully tried to~find evidence of~larger-scale structures (lobes, nebula,\\ldots) around the~object using a~narrow-band H$\\alpha$ filter.  \\cite{zicksl89} assumed that~a~circumstellar disc structure causes a~direction-dependent polarisation of~the~emitted photons. They determined an~intrinsic polarization of~$\\sim$~2\\% from~their observations. On~the~basis of~this value and~the~assumption that~the~scattering plane is~parallel to~the~orbital plane, \\cite{Zickgraf01} estimated an~inclination angle $\\geq 30-45^{\\circ}$.  The gaseous and dusty envelope of~V2028~Cyg makes it very difficult to~determine the~fundamental parameters of the object, since the~commonly used stellar atmosphere models are not applicable in~this case. The~properties of~these objects can only be described by~multidimensional models that combine time-dependent hydrodynamics and~NLTE radiative transfer. This task is our ultimate goal. Our time-dependent hydrodynamical code \\citep{ufo}, which is now extended into 2D, gives the~temperature and~density distributions and~velocity of~the~material for~the~2D radiative transfer code \\citep{osa, disk}. However, this is a~long-term project complicated the~role of~different physical effects  in~the~formation of~the~B[e] phenomenon remaining unclear. Therefore, we decided to~focus our efforts here on~the~analysis of~the~object's temporal behaviour that~may reveal important physical mechanisms and~narrow down the~possible system configurations. The~first papers of~the~planned series focus on~a~description of~the~observed spectral properties of~selected B[e] stars over~a~time interval of~several years.  The~observations, data reduction, and~main spectral features are~described in~Sects. \\ref{obs} and~\\ref{spec}. The following Sect.~\\ref{srv} deals with~the~temporal variability of~the~emissions of~\\ion{Fe}{ii} (6318, 6456, 6417~\\AA), [\\ion{O}{i}]~(6300~\\AA), and~H$\\alpha$. In~Sect. \\ref{dis}, we discuss the~possible nature of~V2028~Cyg based on~our observational results and~give several criteria and~suggestions for~future modelling. ",
        "conclusions": ""
    },
    "1208/1208.1760_arXiv.txt": {
        "abstract": "It is likely that unambiguous habitable zone terrestrial planets of unknown water content will soon be discovered. Water content helps determine surface land fraction, which influences planetary weathering behavior. This is important because the silicate weathering feedback determines the width of the habitable zone in space and time.  Here a low-order model of weathering and climate, useful for gaining qualitative understanding, is developed to examine climate evolution for planets of various land-ocean fractions.  It is pointed out that, if seafloor weathering does not depend directly on surface temperature, there can be no weathering-climate feedback on a waterworld.  This would dramatically narrow the habitable zone of a waterworld. Results from our model indicate that weathering behavior does not depend strongly on land fraction for partially ocean-covered planets.  This is powerful because it suggests that previous habitable zone theory is robust to changes in land fraction, as long as there is some land. Finally, a mechanism is proposed for a waterworld to prevent complete water loss during a moist greenhouse through rapid weathering of exposed continents.  This process is named a ``waterworld self-arrest,'' and it implies that waterworlds can go through a moist greenhouse stage and end up as planets like Earth with partial ocean coverage. This work stresses the importance of surface and geologic effects, in addition to the usual incident stellar flux, for habitability. ",
        "introduction": "\\label{sec:intro}   The habitable zone is traditionally defined as the region around a star where liquid water can exist at the surface of a planet \\citep{Kasting93}. Since climate systems include both positive and negative feedbacks, a planet's surface temperature is non-trivially related to incident stellar flux.  The inner edge of the habitable zone is defined by the ``moist greenhouse,'' which occurs if a planet becomes hot enough (surface temperature of $\\approx$340~K) that large amounts of water can be lost through disassociation by photolysis in the stratosphere and hydrogen escape to space. The outer edge occurs when CO$_2$ reaches a high enough pressure that it can no longer provide warming, either because of increased Rayleigh scattering or because it condenses at the surface, which results in permanent global glaciation.  These limits do not necessarily represent hard limits on life of all types. For example, life survived glaciations that may have been global (``Snowball Earth'') on Earth 600--700 million years ago \\citep{Kirschvink92,Hoffman98}. Furthermore, if greenhouse gases other than CO$_2$ are considered, e.g., hydrogen \\citep{Pierrehumbert:2011p3366,Wordsworth2012-transient}, the outer limit of the habitable zone could be pushed beyond the CO$_2$ condensation limit. Alternative limits on habitability have also been proposed. For example, a reduction in the partial pressure of atmospheric CO$_2$ below $10^{-5}$ bar may prevent C$_4$ photosynthesis, which could curb the complex biosphere \\citep{CALDEIRA:1992p2366,VonBloh:2005p2325}, although it would certainly not limit most types of life.  Studies of planetary habitability have become increasingly pertinent as capabilities for the discovery and characterization of terrestrial or possibly terrestrial exoplanets in or near the habitable zone have improved. For example, planets GJ581d \\citep{Udry:2007p3239,Mayor:2009p3227,Vogt:2010p3247,Wordsworth:2011p3221}, HD85512b \\citep{Pepe11,Kaltenegger11}, Kepler-22b \\citep{Borucki11}, and GJ 677c \\citep{Anglada12} have recently been discovered and determined to be potentially in the habitable zone.  It is likely that the number of good candidates for habitable terrestrial exoplanets will increase in the near future as the Kepler mission, ground-based transit surveys, radial velocity surveys, and gravitational microlensing studies continue to detect new planetary systems.  The ``Faint Young Sun Problem'' discussed in the Earth Science literature is analogous to the habitable zone concept. A planet orbiting a main-sequence star receives ever increasing incident flux (``insolation'') as the star ages. This should tend to hurry a planet through the habitable zone unless the planetary albedo or thermal optical thickness (greenhouse effect) adjust as the system ages.  As pointed out by \\citet{SAGAN:1972p1233}, geological evidence for liquid water early in Earth's history suggests that just such an adjustment must have occurred on Earth. Although there is some debate about the details \\citep{Kasting:2010p1226}, it is fairly widely accepted that the silicate-weathering feedback \\citep{Walker-Hays-Kasting-1981:negative} played an important role in maintaining clement conditions through Earth's history \\citep{Feulner:2012p3570}. Through this negative feedback the temperature-dependent weathering of silicate rocks on continents, which represents the main removal process for CO$_2$ from the atmosphere, is reduced when the temperature decreases, creating a buffering on changes in temperature since CO$_2$ is a strong infrared absorber. The silicate weathering feedback also greatly expands the habitable zone annulus around a star in space \\citep{Kasting93}, since moving a planet further from the star decreases the insolation it receives just as moving backward in time does. If insolation boundaries are used to demarcate the habitable zone, the concept can be used to consider habitability as a function of position relative to the star at a particular time, or as a function of time for a planet at a constant distance from a star.  In addition to continental silicate weathering, weathering can also occur in hydrothermal systems in the basaltic oceanic crust at the seafloor. Seafloor weathering is poorly constrained on Earth; however, it is thought to be weaker than continental weathering and to depend mainly on ocean chemistry, pH, and circulation of seawater through basaltic crust, rather than directly on surface climate \\citep{CALDEIRA:1995p2285,Sleep:2001p3368,Lehir08}. If this is true, CO$_2$ would be less efficiently removed from the atmosphere of a planet with a lower land fraction, leading to higher CO$_2$ levels and a warmer climate.  Furthermore, a planet with a lower land fraction could have a weaker buffering to changes in insolation than a planet with a higher land fraction, which would cause it to have a smaller habitable zone.  While the weathering behavior of a planet could depend on surface land fraction, the water complement of planets in the habitable zone should vary substantially. The reason for this is that the habitable zone is in general located closer to the star than the snow line, the location within a protoplanetary disk outside of which water ice would be present and available to be incorporated into solids. In the solar system, for example, the current habitable zone ranges from approximately 0.8 to 1.7 AU \\citep{Kasting93} if one neglects the uncertain effects of CO$_2$ clouds. If CO$_2$ clouds are assumed to produce a strong warming \\citep{Forget:1997p3442}, the outer limit of the habitable zone could be extended to $\\approx$2.4 AU \\citep{Mischna:2000p3444,Selsis:2007p3441,Kaltenegger:2011p3443}, although recent three-dimensional simulations including atmospheric CO$_2$ condensation suggest that the warming effect may in fact be fairly modest \\citep{Wordsworth:2011p3221,Forget2012,Wordsworth2012}. The snow line is thought to have been around $\\sim$2.5 AU \\citep{morbidelli00}; beyond this distance, evidence for water in the form of ice or hydrated minerals is seen in asteroids and planets.  The general picture for water delivery to Earth has been that the orbits of bodies from beyond the snow line were excited to high eccentricities through gravitational scattering by massive planetary embryos and a young Jupiter.  These high eccentricities would have put water-bearing planetesimals and embryos on Earth-crossing orbits.  A fraction of such bodies would have been accreted by Earth, stochastically delivering volatiles to the young planet \\citep{morbidelli00,obrien06,Raymond:2009p2536}.  Efforts have been made to extend this analysis to planetary systems around other stars.  The same dynamical effects are expected to occur, with the amount of water delivered to a potentially habitable planet being strongly dependent on the presence or orbits of giant planets \\citep{obrien06,raymond06,Raymond:2009p2536}.  As low mass stars are expected to have lower mass disks and therefore fewer massive bodies to produce gravitational scattering, low mass stars may be more likely to have volatile-poor habitable-zone planets \\citep{Raymond:2007}. Stars of mass $\\sim$1~$M_{\\odot}$ or higher, on the other hand, could easily have habitable-zone planets of similar or greater fractional mass of water than Earth \\citep{Raymond:2007}.  As described above, current thinking mostly focuses on hydrated asteroids as the main source of water for a habitable planet; however, there are other dynamical mechanisms that could allow habitable planets to accrete significant amounts of water and volatiles. Cometary-bodies would deliver a larger amount of water for a given impactor mass than asteroids.  Should such bodies get scattered over the course of planet formation, this may allow larger amounts of water to be accreted than otherwise predicted.  Furthermore, \\citet{kutchner03} argued that planets which formed outside the snow line could migrate inwards due to gravitational torques from a protoplanetary disk or scattering with other planets.  Such bodies would naturally accrete large fractions of water ice during formation beyond the snow line, which would then melt and sublimate in a warmer orbit, providing a high volatile content to a planet close to its star.  Habitable zone planets could therefore have a wide variety of water mass fractions, which would lead to varying land fractions. Furthermore, even for a constant mass fraction of water, scaling relations dictate that land fraction depends on planetary size. At the same time, the surface land fraction could exert strong control on the planetary carbon cycle, which strongly influences planetary habitability. This warrants a general consideration of weathering behavior on planets of varying land fraction.  The main objective of this paper is to investigate the effect of land fraction on the carbon cycle and weathering behavior of a terrestrial planet in the habitable zone. We will outline and use a simple analytical model for weathering and global climate that necessarily makes grave approximations to the real physical processes. For example, we will use existing parameterizations of seafloor weathering, while acknowledging that observational and experimental constraints on such parameterizations are minimal. We will only use the model, however, to make statements that do not depend strongly on uncertain aspects of the parameterizations. This model should be used to understand intuitively the qualitative behavior of the system rather than to make quantitative estimates. A major strength of the model is that it is easy to derive and understand, yet should capture the most significant physical processes. This type of modeling is appropriate in the study of exoplanets, for which limited data that would be relevant for a geochemical model exist.  We will consider an Earth-like planet with silicate rocks, a large reservoir of carbon in carbonate rocks, and at least some surface ocean. We will use equilibrium relations for weathering, which is reasonable for the slow changes in insolation that a main-sequence star experiences. The climate model we use is a linearization of a zero dimensional model. Although this is a severe approximation, it allows the analytical progress that we feel is useful for obtaining insight into the problem. We will consider planetary surface land fractions ranging from partial ocean coverage to complete waterworlds. In this context a planet would be a waterworld if the highest land were covered by even 1~m of water, although a planet with more water than this would qualify as a waterworld as well. However we will assume the geophysical context is a planet with a substantial rock mantle and only up to roughly $\\approx$10 times Earth's water mass fraction \\citep[0.02\\%-0.1\\%;][]{Hirschmann:2012p3481} rather than potential waterworlds that might be \\cal{O}(10\\%) or more water by mass \\citep{Fu2010} and would therefore have vastly different volatile cycles. Although partially ocean-covered planets have recently been referred to as ``aquaplanets'' \\citep{abe2011}, we will not adopt this terminology because ``aquaplanet'' has a long history of being used to denote a completely ocean-covered planet in the climate and atmospheric dynamics communities.   We will find that the weathering behavior is fairly insensitive to land fraction when there is partial ocean coverage. For example, we will find that weathering feedbacks function similarly, yielding a habitable zone of similar width, if a planet has a land fraction of 0.3 (like modern Earth) or 0.01 (equivalent to the combined size of Greenland and Mexico). In contrast, we will find that the weathering behavior of a waterworld is drastically different from a planet with partial ocean coverage. If seafloor weathering depends mainly on ocean acidity, rather than planetary surface temperature, no weathering feedback operates on a waterworld and it should have a narrow habitable zone and progress through it quickly as its star ages. Finally, we will argue that it is possible for a waterworld to stop a moist greenhouse in progress when continent is exposed by drawing down the CO$_2$ through massive weathering, which would leave the planet in a clement state with partial ocean coverage. We will refer to this possibility as a ``waterworld self-arrest.''  This paper complements in two ways the recent work of \\citet{abe2011}, who found that a nearly dry planet should have a wider habitable zone than a planet with some water. First, the calculations made by \\citet{abe2011} focused on climate modeling rather than weathering, although they included a qualitative description of factors that would influence weathering on a dry planet. Second, we consider a variety of land fractions up to the limiting case of a waterworld.  The outline of this paper is as follows. We describe our model in Section \\ref{sec:model}, use it in Section \\ref{sec:results}, and perform a sensitivity analysis in Section \\ref{sec:sensitivity}. We discuss the possibility of a waterworld self-arrest in Section \\ref{sec:self-arrest}. We outline observational implications and prospects for confirmation or falsification of our work in Section \\ref{sec:observables}. We discuss our results further in Section \\ref{sec:discussion}, including considering the limitations of our various assumptions, and conclude in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that using standard weathering parameterizations, the weathering behavior of a partially ocean-covered Earth-like planet does not depend strongly on land fraction, as long as the land fraction is greater than $\\approx$0.01. Consequently, planets with some continent and some ocean should have a habitable zone of similar width. This is a powerful result because it indicates that previous habitable zone theory developed assuming Earth-like land fraction and weathering behavior should be broadly applicable.  We have also pointed out that, as long as seafloor weathering does not depend directly on surface temperature, a climate-weathering feedback cannot operate on a waterworld. This is a significant result because it would imply that waterworlds have a much narrower habitable zone than a planet with even a few small continents. We find, however, that weathering could operate quickly enough that a waterworld could ``self-arrest'' while undergoing a moist greenhouse and the planet would be left with partial ocean coverage and a clement climate. If this result holds up to more detailed kinetic weathering modeling, it would be profound, because it implies that waterworlds that form in the habitable zone have a pathway to evolve into a planet with partial ocean coverage that is more resistant to changes in stellar luminosity."
    },
    "1208/1208.3079_arXiv.txt": {
        "abstract": "We present a pan-chromatic analysis of an unprecedented sample of 1402 250\\,$\\mu$m-selected galaxies at $z < 0.5$ ($\\bar z = 0.24$) from the {\\it Herschel}-ATLAS survey. We complement our {\\em Herschel\\/} 100--500\\,$\\mu$m data with UV--K-band photometry from the Galaxy And Mass Assembly (GAMA) survey and apply the {\\sc magphys} energy-balance technique to produce pan-chromatic SEDs for a representative sample of 250\\,$\\mu$m selected galaxies spanning the most recent 5 Gyr of cosmic history. We derive estimates of physical parameters, including star formation rates, stellar masses, dust masses and infrared luminosities. The typical H-ATLAS galaxy at $z<0.5$ has a far-infrared luminosity in the range $10^{10} - 10^{12}~L_\\odot$ (SFR: 1--50 $M_{\\odot}~yr^{-1}$) thus is broadly representative of normal star forming galaxies over this redshift range. We show that 250\\,$\\mu$m selected galaxies contain a larger mass of dust at a given infra-red luminosity or star-formation rate than previous samples selected at 60\\,$\\mu$m from {\\em IRAS\\/}. We derive typical SEDs for \\hatlas\\ galaxies, and show that the emergent SED shape is most sensitive to specific star-formation rate. The optical-UV SEDs also become more reddened due to dust at higher redshifts. Our template SEDs are significantly cooler than existing infra-red templates. They may therefore be most appropriate for inferring total IR luminosities from moderate redshift submillimetre selected samples and for inclusion in models of the lower redshift submillimetre galaxy populations. ",
        "introduction": "In the past couple of decades, our understanding of the Universe has flourished as a result of our new-found ability to observe in almost all regions of the electromagnetic spectrum. This advance is in no small part due to our ability to associate observations at different wavelengths with particular astrophysical phenomena and link them together, making modern astronomy truly pan-chromatic. By observing an astronomical source at multiple wavelengths, we may piece together its spectral energy distribution (SED), and by comparing the observed SED to models, we may infer the physical properties of the source (or sample of sources) that we are studying.  At ultraviolet, optical and near-infrared wavelengths the SED of the average galaxy is dominated by emission from stars; there are many tens of different models to which we may compare our observations, in the hope of understanding the stellar components of astrophysical sources (e.g. Bruzual \\& Charlot, 2003, Fioc \\& Rocca-Volmerange, 1997, V\\'azquez \\& Leitherer, 2005, Anders \\& Fritze-von Alvensleben, 2003, Jimenez et al. 1995, 2004, Maraston, 2005, Pietrinferni et al. 2004). Such SED model analysis may be used to determine basic properties of a galaxy's stellar components, such as its age, metallicity or stellar mass (see e.g. Carter et al., 2009, Smith \\& Jarvis, 2007, Collins et al., 2009, Pacifici et al., 2012, Pforr, Maraston \\&\\ Tonini, 2012).  While the ultraviolet to near-infrared emission tells us about the stellar content of a normal galaxy (subject to correcting for attenuation by dust of the different stellar components; Charlot \\&\\ Fall, 2000, Tuffs et al. 2004, Pierini et al. 2004), the far-infrared and sub-millimetre wavelengths probe its cool dust content, which is itself crucial to our understanding of star formation, since approximately half of the energy ever radiated by stars has been absorbed by dust and re-radiated at these wavelengths (e.g. Puget et al., 1996, Fixsen et al., 1998). The sub-millimetre region has been a difficult part of the electromagnetic spectrum in which to conduct large galaxy surveys (e.g. Smail et al. 1997, Hughes et al. 1998, Eales et al. 1999). Previous ground-based sub-millimetre surveys had to be either pointed at pre-selected targets, or limited to relatively small regions of sky covering areas of $<1$ deg$^2$ (Coppin et al., 2006, Weiss et al., 2009). The combined effects of the large negative $k$-correction at these wavelengths, sensitivity and the steep number counts, have meant that the average 850\\,$\\mu$m selected sub-millimetre galaxy is extremely luminous ($10^{12}-10^{13}~L_\\odot$) and resides at high redshift ($z \\sim 2$, e.g. Chapman et al. 2005). Few relatively local galaxies have been found in blind sub-mm surveys, due to the small local volumes probed in these surveys coupled with the observing wavelength targetting the faint Rayleigh-Jeans tail of the dust SED at low redshift. Our undertstanding of the local Universe at sub-mm wavelengths has come so far from targetted surveys such as the SCUBA Local Universe Galaxy Survey (SLUGS, Dunne et al. 2000), which observed a sample of 184 \\iras- and optically-selected galaxies (Vlahakis, Dunne \\& Eales, 2005). Pre-selected galaxies in this way can lead to biases if there are classes of sub-mm emitting galaxies which are not bright at either optical or 60\\,$mu$m wavelengths. The SLUGS surveys were also limited to very nearby galaxies and so could not address the question of evolution of sub-mm properties in the relatively recent past.  With the advent of the PACS (Poglitsch et al. 2010) and SPIRE (Griffin et al. 2010) instruments aboard the ESA {\\it Herschel Space Observatory} (Pilbratt et al. 2010), we now have our first opportunity to survey a large area of sky at sub-mm wavelengths. The angular resolution and sensitivity of {\\it Herschel} allow us to robustly determine the counterparts to thousands of local sub-millimetre {\\it selected} galaxies across the whole electromagnetic spectrum, thus gaining invaluable insight into their physical processes. This paper uses a sample from the {\\em Herschel} Astrophysical TeraHertz Large Area Survey ({\\it H}-ATLAS: Eales et al. 2010) and presents fits to their UV--sub-mm SEDs.  This is the first relatively local ($z<0.5$) sub-mm selected sample for which such complete SED modelling has been performed. This work is based on only 3 percent of the final data-set, but is still large enough to provide a statistical study of the optical and IR properties of $>1000$ 250\\,$\\mu$m selected galaxies, and templates for SEDs which can be applied more widely.  Studies of the multi-wavelength properties of the relatively small number of galaxies detected in sub-millimetre surveys have been extensive (e.g. Swinbank et al. 2009). At high redshifts, galaxy star formation rates have been frequently estimated based on a single sub-millimetre flux measurement (e.g. at 850\\,$\\mu$m), and a local template SED belonging to e.g. M82 or Arp\\,220 (e.g. Silva et al., 1998), chosen not because they are known to be representative of the average sub-millimetre galaxy, but rather because they are comparatively well-studied.  Another commonly-used method of describing far-infrared galaxy SEDs is to assume one or more components with modified black-body (the so-called ``grey-body'') profiles. In these simple parametrisations, the observed flux densities depend only on the temperatures ($T$) and dust emissivity index ($\\beta$), which may be either assumed or derived, depending on the available observations. Such simple grey-body profiles have been widely shown to broadly reproduce the sparsely-sampled far-IR SEDs of galaxies at all redshifts (e.g. Dunne et al. 2000, Blain et al. 2002, Blain, Barnard \\&\\ Chapman 2003, Pope et al. 2006, K\\'ovacs et al. 2006, Dye et al., 2010), although when the SED is sampled from 60$\\mu$m to the sub-mm additional greybody components may be required to reproduce the observations (e.g. Dunne \\&\\ Eales, 2001, Galametz et al., 2011, Smith et al. 2010, Dale et al., 2012).   Empirical templates have been created for use with sparsely-sampled far-infrared data based on observations of small samples of local galaxies with good coverage from mid- to far-IR wavelengths (e.g. Chary \\&\\ Elbaz, 2001, Dale \\&\\ Helou, 2002, Rieke et al. 2009). Selecting galaxies at shorter FIR wavelengths tends to favour those with substantial warm dust components, which may not be representative of populations selected at longer wavelengths with {\\em Herschel\\/} and ground based sub-mm instruments. Several studies have found that sub-mm selected galaxies (so far mostly at higher redshifts) may have colder dust than their local equivalents at similar far infrared luminosities (e.g. Pope et al. 2006, Coppin et al. 2008, Hwang et al., 2010, Magnelli et al., 2012).  In this paper, we use a model that relies on energy balance - the idea that the energy absorbed by dust at ultra-violet and optical wavelengths must be re-radiated in the far-infrared - combined with a statistical fitting approach, to consistently model each galaxy's full SED, and gain robust constraints on the star formation activity, stellar and dust content of 250\\,$\\mu$m selected galaxies from \\hatlas.  In Section \\ref{sec:data} we discuss the {\\it Herschel}-ATLAS survey and the multi-wavelength data used in generating the catalogue, while in Section \\ref{sec:method} we discuss the SED-fitting method used in the analyses which we present in Section \\ref{sec:Results}. In Section \\ref{comparisons} we compare our median SEDs with other templates available, and in Section \\ref{conclusions} we present some conclusions based on our results for the population of sub-millimetre galaxies in general. Throughout this paper, we use a standard cosmology with $H_0 = 71$\\,km\\,s$^{-1}$\\,Mpc\\,$^{-1}$, $\\Omega_M = 0.27$ and $\\Omega_\\Lambda = 0.73$, and an Initial Mass Function (IMF) from Chabrier (2003). ",
        "conclusions": "\\label{conclusions}  We have determined SEDs for a total of 1402 250$\\mu$m-selected galaxies from the Herschel-ATLAS science demonstation catalogue with reliable counterparts and matched aperture photometry from the $u$- to $K$-bands from the GAMA database. We also include far and near ultraviolet data from the {\\it GALEX}-GAMA survey, as well as the {\\it H}-ATLAS data from PACS and SPIRE. Of these 1402 galaxies, 1289 are well described by the model of DCE08, and we use these SEDs and the model parameter probability density functions derived from the energy balance SED-fitting, to determine the properties of these 250\\,$\\mu$m selected galaxies out to $z=0.5$. \\begin{enumerate} \\item Studies of the colours of galaxies in our sample, and a suite of simulations, suggest that our sample is representative of the broader population of 250\\,$\\mu$m galaxies out to $z < 0.35$. \\item The average \\hatlas\\ galaxy in our sample has a star formation rate of $\\sim$\\,4.0~M$_\\odot $yr$^{-1}$, L$_{dust} \\approx 6.4 \\times 10^{10}$ ~L$_\\odot$, and a dust to stellar mass ratio of $\\sim$0.4 per cent, while the median redshift is $z=0.24$. \\item Our results support the idea that \\iras\\ and \\hatlas\\ selected galaxies in the local Universe are different populations. Due to its lack of sensitiity and short selection wavelength, \\iras\\ preferentially selected galaxies with larger warm dust content, and consequently these galaxies are more luminous in the infrared for a given mass of dust. The \\hatlas\\ selection at 250\\,$\\mu$m is less biased towards strongly star forming objects over the same redshift range because of the longer selection wavelength and far superior sensitivity compared to \\iras. \\iras\\ misses a population of massive dusty galaxies with colder dust temperatures, as was shown previously by Vlahakis, Dunne \\& Eales (2005). \\item The correlation between star formation rate and dust mass presented in da Cunha et al. (2010) is also present in this sample, although {\\it Herschel} ATLAS-selected galaxies contain larger dust masses for a given star formation rate compared to the \\iras\\ selected sample of dC10. There is also a correlation between specific dust mass (\\mdust\\slash\\mstar) and SSFR, which is not well reproduced by simple chemical and dust evolution models. \\item The specific star formation rate of lower mass galaxies ($\\log_{10} M_{\\mathrm{star}}\\slash \\Msolar < 10.2$) is higher than that of the most massive galaxies in our sample (those with $\\log_{10} M_{\\mathrm{star}}\\slash \\Msolar > 10.6$) at all redshifts, supporting previous results that lower mass galaxies dominate the star formation rate density in the local universe. \\item Stacks of SEDs show that sSFR is the strongest galaxy property driving the SED shape across both the UV/optical and FIR, as first noticed in the smaller sample of DCE08. Trends with \\ldust\\ are much weaker since smaller mass galaxies will have low \\ldust\\ and yet could have the highest values of sSFR. We see a signficant trend in this sample for galaxies to have more obscured optical/UV SEDs and higher reprocessed fractions with increasing redshift. \\item Existing templates for panchromatic SEDs of galaxies show shorter FIR peaks and excess mid-IR emission compared to median stacked SEDs of galaxies in our \\hatlas\\ sample (binned by \\ldust) although the mid-IR discrepancy is not significant at this time due to our lack of mid-IR data to constrain this part of the SED. Templates from Rieke et al. (2009) are the closest match to ours in terms of the FIR properties although they still predict a warmer FIR peak at the highest luminosities compared to our findings. We provide a new set of panchromatic SED templates from the UV--sub-mm, to enable more representative studies of dusty galaxies in the local Universe in the {\\it Herschel} era. \\item Data from the \\wise\\ satellite, which covers the wavelength range between 3 and 23\\,$\\mu$m will provide valuable constraints to the mid-IR and PAH features, as well as the hot dust component of these local galaxies. It will be interesting to see if the differences between templates in the mid-IR region persists when these data are included in the fitting. \\end{enumerate}"
    },
    "1208/1208.1759.txt": {
        "abstract": "We present resolved \\emph{Herschel} images of circumbinary debris disks in the \\alp~(HD 139006) and \\bet~(HD13161) systems. By modelling their structure, we find that both disks are consistent with being aligned with the binary orbital planes. Though secular perturbations from the binary can bring the disk into alignment, in both cases the alignment time at the distances at which the disk is resolved is greater than the stellar age, so we conclude that the coplanarity was primordial. Neither disk can be modelled as a narrow ring, requiring extended radial distributions. To satisfy both the \\emph{Herschel} and mid-IR images of the \\alp~disk, we construct a model that extends from 1-300AU, whose radial profile is broadly consistent with a picture where planetesimal collisions are excited by secular perturbations from the binary. However, this model is also consistent with stirring by other mechanisms, such as the formation of Pluto-sized objects. The \\bet~disk is modelled as a disk that extends from 50-400AU. A model with depleted (rather than empty) inner regions also reproduces the observations and is consistent with binary and other stirring mechanisms. As part of the modelling process, we find that the \\emph{Herschel} PACS beam varies by as much as 10\\% at 70$\\mu$m and a few \\% at 100$\\mu$m. The 70$\\mu$m variation can therefore hinder image interpretation, particularly for poorly resolved objects. The number of systems in which circumbinary debris disk orientations have been compared with the binary plane is now four. More systems are needed, but a picture in which disks around very close binaries (\\alp, \\bet, and HD 98800, with periods of a few weeks to a year) are aligned, and disks around wider binaries (99 Her, with a period of 50 years) are misaligned, may be emerging. This picture is qualitatively consistent with the expectation that the protoplanetary disks from which the debris emerged are more likely to be aligned if their binaries have shorter periods. ",
        "introduction": "\\label{s:intro}  The \\emph{Herschel} Key Program DEBRIS (Dust Emission via a Bias free Reconnaissance in the Infrared/Submillimeter) has observed a large sample of nearby stars to discover and characterise extrasolar analogues to the Solar System's asteroid and Edgeworth-Kuiper belts, collectively known as ``debris disks.'' The 3.5m \\emph{Herschel} mirror diameter provides 6-7'' resolution at 70-100$\\mu$m \\citep{2010A&A...518L...1P}, and as a consequence our survey has resolved many disks around stars in the Solar neighbourhood for the first time \\citep[][Wyatt et. al. 2012]{2010A&A...518L.135M,2011MNRAS.417.1715C,2012MNRAS.421.2264K}.\\footnote{\\emph{Herschel} is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.}  Here we present resolved images of circumbinary disks in the \\alp~and \\bet~systems. These systems are interesting because unlike most debris disk+binary systems, the binary orbits are well characterised. The combination of a known orbit and resolved disk means we can compare their relative inclination.  Our observations of a disk around the binary 99 Her \\citep{2012MNRAS.421.2264K} were a step toward building on the binary debris disk study of \\citet{2007ApJ...658.1289T}. Their \\emph{Spitzer} study found that debris disks are generally as common in binary systems as in single systems, but are less likely to reside in systems with binary separations in the 3-30AU range \\citep[see also][]{2012ApJ...745..147R}. However, only some of their systems had detections at multiple wavelengths to constrain the disk location and none were reported as resolved, making the true dust location uncertain. Even in the case of dust detection at multiple wavelengths, the true dust location is unknown because grains of different compositions and sizes can have the same temperature at different stellocentric distances. In addition to uncertainty in the dust location, only the projected sky separation of the binary (not the binary semi-major axis) was generally known, adding further uncertainty.  Systems with resolved disks and well characterised binary orbits, such as 99 Her, \\alp, \\bet, and HD 98800 \\citep{2010ApJ...710..462A} remove these ambiguities. The dust location and structure can be inferred within the context of the binary orbit, leading to robust conclusions about whether the dust resides on stable orbits. One can also consider whether perturbations from the binary play an important role in setting the dust dynamics. For example, in the 99 Her system the disk position angle appears misaligned with the binary line of nodes. This misalignment is best explained with particles on polar orbits because these are stable for the stellar lifetime \\citep{2012MNRAS.421.2264K}. Another question is whether binary perturbations can ``stir'' the disk particles by increasing their inclinations and eccentricities, eventually resulting in high enough relative velocities that collisions are destructive. This process is analogous to the planet stirring model proposed by \\citet{2009MNRAS.399.1403M}, though may rely on vertical (inclination) stirring because companions of comparable mass induce lower forced eccentricities than companions that are much less massive than the star \\citep{2012MNRAS.421.2264K,2004ApJ...609.1065M}.  An additional link can be made to star and planet formation. Young binary systems with small to medium ($\\lesssim$100AU) separations are thought to form coplanar with their protoplanetary disks, because the disk torque aligns the binary orbit on timescales short relative to the disk lifetime \\citep[e.g.][]{2000MNRAS.317..773B}. Testing this prediction is difficult because it is difficult to ascertain disk and binary orientations at the distances of the nearest star forming regions \\citep[e.g.][]{2007prpl.conf..395M}. Because debris disks around older main-sequence binaries should retain the same orientation as the protoplanetary disks from which they emerged, these disks yield information on the outcome of star formation. The advantage is that compared to star-forming regions, these systems are much closer to Earth and hence larger on the sky and brighter, making disk and binary characterisation much easier.  Though planet formation may be hindered to some degree by high collision velocities induced by binary perturbations \\citep[e.g.][]{2004ApJ...609.1065M,2007MNRAS.380.1119S}, the discovery of several circumbinary planets shows that planets do indeed form around binary stars \\citep[e.g.][]{2011Sci...333.1602D,2012Natur.481..475W}. Few such systems are known because binaries are generally avoided by radial velocity surveys \\citep[see][]{2005ApJ...626..431K,2009ApJ...704..513K}. However, the very existence of circumbinary debris disks provides evidence that planet formation around binaries can proceed to form at least 10-100km sized objects, which must exist to feed the observed dust through collisions \\citep[e.g.][]{2008ARA&A..46..339W,2010RAA....10..383K}. Further, circumbinary disks are relatively common \\citep{2007ApJ...658.1289T}, suggesting that circumbinary planets may be no less unusual than their circumstellar equivalents.  This paper is laid out as follows. We first consider the stellar and orbital properties, along with previous IR observations of the \\alp~and \\bet~systems (\\S \\ref{s:stars}). We then show the \\emph{Herschel} data (\\S \\ref{s:obs}) and simple models (\\S \\ref{s:mod}), and then interpret these within the context of the expected dynamics (\\S \\ref{s:dyn}). We discuss the results and conclude in \\S\\S \\ref{s:disc} \\& \\ref{s:sum}. ",
        "conclusions": "\\label{s:disc}  By considering the resolved models and expected dynamics, we have shown that the circumbinary debris disks in both the \\alp~and \\bet~systems probably formed coplanar with their parent binaries. A corollary is that because the disks were primordially aligned, there should be little vertical structure inside the alignment radius. Both disks were successfully modelled as disks with a single plane of symmetry, but are too poorly resolved with \\emph{Herschel} to strongly verify this statement. However, for \\alp~the low-level contours from the 11 and 18$\\mu$m mid-IR imaging suggest that the inner disk has a similar position angle to the binary and outer disk. Only one of the 18$\\mu$m images shows the same position angle, though the fact that the other is actually narrower than the PSF in the same direction makes this extension questionable, as discussed by \\citet{2010ApJ...723.1418M}.  While there appears to be no need or evidence for vertical disk structure induced by perturbations to disk particles' inclinations, an indirect signature could exist from inclination and eccentricity variations imposed on disk particles. These variations ``stir'' the disk, where we are defining ``stirring'' to be any mechanism that increases relative velocities between particles sufficiently that collisions become catastrophically destructive, thus initiating a collisional cascade. In stirred regions mass at the small end of the collisional cascade is removed by radiation forces, thus depleting the disk. In unstirred regions the mass remains constant. The size distribution in stirred regions has many more small grains for a given mass in large objects, so is more easily detected due to the larger emitting surface area.\\footnote{Once catastrophic collisions are occurring, the level of stirring is also important. An increase in stirring does not have a significant effect on the collision rate between particles of similar sizes because higher velocities are accompanied by a lower volume density of particles (e.g. the scale height and/or radial disk extent increases with the velocities). However, particles are typically destroyed in a collision with a much smaller particle, whose size is set by the collision velocity. Thus, higher levels of stirring mean that smaller particles are capable of destroying objects. In standard collisional size distributions \\citep[e.g.][]{1969JGR....74.2531D} the number of particles, and therefore the number of potential destructive impactors, increases strongly with decreasing size, so the disk depletion rate due to collisions increases with increased stirring.}  Though the structure of the inner disk is not well constrained, we have shown that a plausible model that satisfies both the \\emph{Herschel} and mid-IR data for \\alp~has a continuous optical depth profile that peaks around 50AU. Because the typical expectation from Solar System and protoplanetary disk studies is for the surface density to decrease with distance from the star \\citep[e.g.][]{1977Ap&SS..51..153W,2009ApJ...700.1502A}, this profile can be interpreted as a depletion of material inside 50AU. Such a depletion is expected in standard models of collisional evolution, where the disk decay rate is faster at smaller semi-major axes \\citep[e.g.][]{2003ApJ...598..626D,2007ApJ...658..569W}. This 50AU turn-over distance is similar to $r_{\\rm align}$ shown in Figure \\ref{fig:sec}, so it is therefore plausible that the depletion inside here is due to increased velocities imposed by binary perturbations.  Considering this binary stirring picture in more detail, vertical (inclination) stirring is probably more important than radial (eccentricity) stirring in circumbinary disks, particularly when the disk lies at large distances relative to the binary separation. As the binary mass ratio decreases (i.e. the stars' masses become more similar), the eccentricities imposed on an exterior particle decrease. For example, using the expression derived by \\citet{2004ApJ...609.1065M}, the ``forced'' eccentricity of particles at 50AU around \\alp~is about 0.001, at least an order of magnitude lower than the eccentricities thought to exist in observed debris disks \\citep[e.g.][]{2004AJ....127..513K,2008ApJ...687..608K}. In contrast, for inclination stirring to be effective the initial disk plane would only need to be a few degrees different to the binary plane (though some non-zero eccentricity is needed to ensure crossing orbits). For example, with an initial misalignment of 1$^\\circ$, collision velocities are stirred to about 0.02 times the Keplerian velocity (150m s$^{-1}$ at 50AU), so objects with dispersal thresholds less than roughly $10^4$J kg$^{-1}$ ($\\lesssim$10km) will be disrupted and dispersed in collisions with similar sized objects. Therefore, the disk decays from the inside-out as secular perturbations increase the relative velocities in particle collisions by aligning the disk, and do so on a shorter timescale closer to the central star(s). Assuming that the initial misalignment was large enough, collisions will be destructive inside some radius that will be close to $r_{\\rm align}$ \\citep[but not exactly at $r_{\\rm align}$,][]{2009MNRAS.399.1403M}. Outside this radius, collisions are unaffected by the presence of the binary and therefore not destructive, providing a possible explanation for the drop in optical depth in our model beyond 50AU.  % print,(4e-4*[50.,100]^3/[3.5,4.9]^1.5/[350.,730.])^2 % print,2.2e-10*[50.,100]^(23/6.)*sqrt(60)*74e3/[3.5,4.9]^(7/3.)/[350.,730.]  While the \\alp~disk is consistent with being stirred by the binary, we also consider two alternative stirring mechanisms. The first is ``pre-stirring'' in which objects are assumed to have been stirred at some very early time by an unspecified mechanism that does not necessarily still operate (e.g. a stellar flyby or the result of gas disk dispersal). A potential issue with such a scenario is that collisional damping may reduce the velocities sufficiently that collisions are not catastrophic for a significant fraction of the stellar main-sequence lifetime \\citep[e.g.][]{2002AJ....123.1757K,2004ARA&A..42..549G}. Whether collisional damping is important depends on the relative sizes of the objects and their destructive impactors, which in turn depends on the relative velocities. If pre-stirred objects can reach sufficiently high eccentricities ($\\gtrsim$0.1, though the value depends on object strength and stellocentric distance), then the disk can remain stirred for the stellar lifetime \\citep{2011ApJ...739...36S}. Such a disk is depleted radially from the inside-out, again because the depletion rate is a strong function of radius. The type of structure that is expected is therefore a radially increasing optical depth profile, which turns over and decreases where collisions have yet to deplete the disk significantly \\citep[e.g.][]{2010MNRAS.405.1253K}. Using equation (6) from \\citet{2010MNRAS.405.1253K} we find for the best fit planetesimal properties from \\citep{2007ApJ...663..365W} that the \\alp~disk would be depleted within 50AU if the disk was stirred from an arbitrarily early time (e.g. pre-stirred) and was about 5-15 times less massive than the solid component of the minimum-mass Solar nebula \\citep[MMSN,][when scaled linearly with the binary mass]{1977Ap&SS..51..153W}.\\footnote{We are not necessarily suggesting that the solid mass in the primordial protoplanetary disk was similarly depleted relative to the MMSN, as the mass could for example have gone into building planets. However, given the large observed dispersion in protoplanetary disk masses in star forming regions, such a depletion is easily possible \\citep[e.g.][]{2005ApJ...631.1134A}}  The second mechanism is ``self-stirring'', where random velocities are initially slow enough that collisions result in accretion and growth. Once the largest objects reach roughly Pluto-size, they increase the velocities of smaller planetesimals and initiate a collisional cascade, and the production of visible levels of dust \\citep{2004AJ....127..513K}. Again, this process works in a radially inside-out fashion, so a self-stirred disk looks similar to a pre-stirred one, with one key difference. Because planetesimals have not been stirred outside where Pluto-sized objects have formed, collisions do not result in a collisional cascade and the disk should show a drop in optical depth beyond this distance \\citep{2010MNRAS.405.1253K}. Using equation (9) from \\citet{2010MNRAS.405.1253K}, for Pluto-sized objects to stir the disk to only 50AU by 350Myr, the disk would have to be 2000 times less massive than a scaled MMSN \\citep[see also][from which the Pluto-formation and stirring times were derived]{2008ApJS..179..451K}. A disk with such low mass would not be visible in a self-stirring scenario \\citep[e.g.][]{2008ApJS..179..451K}, which appears to disfavour this scenario. However, this calculation assumed that the planetesimals started out with 1m-1km sizes. If the planetesimals were initially much larger, the time to form Pluto-sized objects would also be longer \\citep[e.g.][]{2010ApJS..188..242K}. This longer formation time would mean that a more massive disk, which would be correspondingly brighter and therefore more consistent with the observations, could form Pluto-sized objects and stir the disk only as far as 50AU by 350Myr.  Therefore all stirring models appear consistent with the observed peak in surface density at 50AU. However, the model derived in \\S \\ref{ss:img} required a drop in surface density beyond 50AU, a feature expected in self-stirred and binary-stirred disks, which would appear to disfavour a pre-stirred interpretation, but not discern between self and binary-stirring. A caveat on this conclusion is that the disk structure is poorly constrained by the low resolution of the observations, and the dust observed with \\emph{Herschel} may not trace the parent body locations, particularly in the outer regions. For example, the peak at 50AU may simply represent the outer edge of a parent body disk that is pre-stirred, and the (decreased) emission beyond 50AU could arise from small grains originating at 50AU forced onto larger eccentric orbits by radiation pressure \\citep[e.g.][]{2003A&A...408..775T,2006A&A...455..509K}.  Regardless of the stirring mechanism, we can compare the disk structures with those expected if they decayed from some arbitrarily large level. In this picture the face-on geometrical optical depth (with the same assumptions used above) is \\begin{equation}\\label{eq:tau} \\tau = 8.9 \\times 10^{-5} r^{7/3} M_\\star^{-5/6} L_\\star^{-0.5} t^{-1} \\end{equation} where $M_\\star$ is the stellar mass in Solar units and $t$ is the system age in Myr \\citep[][equation 8]{2010MNRAS.405.1253K}. The $r^{7/3}$ dependence gives the expected radial profile of a disk whose planetesimal properties are the same everywhere. With these assumptions, for \\alp~the expected optical depth at 350Myr is $1.15 \\times 10^{-8}$ at 1AU, which is somewhat smaller than the model value of $6.4 \\times 10^{-8}$. Given that there is considerable uncertainty in the planetesimal properties and the true dust distribution, we do not consider this difference a cause for concern. The model does not increase as strongly with radius as Equation (\\ref{eq:tau}), so is below the expected level outside a few AU anyway. That the disk model has a $r^{1.7}$ dependence rather than $r^{7/3}$ could indicate that the planetesimal properties have a radial dependence, for example that they become smaller or weaker at larger distances (the model parameters are very uncertain however, so such a dependence is not required).  The \\bet~data are consistent with a continuous optical depth profile that peaks around 100AU.  Given that the secular precession time depends strongly on semi-major axis, and that the stellar age is uncertain, this distance is not significantly inside $r_{\\rm align}$ and the disk is therefore plausibly binary stirred. The expected optical depth at 730Myr is $3.8 \\times 10^{-9}$ at 1AU, which is smaller than the model value of $2.0 \\times 10^{-8}$. Again, the discrepancy is not particularly large given the model assumptions and uncertainty. Using the same equations as above, the \\bet~disk would be depleted out to 100AU by 730Myr for a disk 1-5 times less massive than a scaled MMSN if it were stirred from the outset. Stirring by Pluto-formation out to this distance only requires a disk 400 times less massive than a scaled MMSN (but again could be more massive if planetesimals are larger). Thus, the disk could be depleted out to 100AU by collisional evolution, but the stirring mechanism is unclear. Unlike \\alp, which has a mid-IR detection, there is no evidence for warm emission in the \\bet~system. Such a detection could break the degeneracy in our models, which cannot tell whether the regions inside 50-100AU are devoid of debris (e.g. due to dynamical clearing by planets), or simply depleted by collisional evolution.  This ``standard'' picture of debris disk stirring and evolution is not the only possibility. For example, catastrophic collisions may be caused due to crossing orbits in a disk where self-gravity is important, with the additional possibility that such disks may appear non-axisymmetric \\citep{2012MNRAS.421.2368J}. It is also possible that some observed debris disks are not stirred to catastrophic collision velocities at all. \\citet{2010MNRAS.401..867H} show that long-lived ``warm'' planetesimal disks could exist, in which collisions are not typically disruptive. They suggest that a test for such a scenario is that the disk spectrum should look like a blackbody, because the disk particles required for such warm disks to survive are large enough that they act like blackbodies (i.e. absorb and emit light efficiently). For the two systems considered here the disk spectra appear to rule out such a scenario because they lie significantly below pure blackbodies beyond wavelengths of about 100$\\mu$m (Fig. \\ref{fig:sed}), suggesting that the particles are emitting inefficiently at long wavelengths and are predominantly smaller than $\\sim$1mm. However, we do not exclude the possibility that the disks are ``warm'', because the small grains that are observed could have been created in erosive and bouncing collisions. Modelling the size distribution would make more quantitative predictions to test this possibility.  Though both disks have extended or continuous dust distributions as observed by \\emph{Herschel}, the parent bodies may (or may not) be more localised. All the models we considered find that the dust optical depth decreases with distance in the outer regions. Qualitatively, this structure is expected when the parent bodies occupy a narrow ``birth'' ring and small grains are placed on eccentric and hyperbolic orbits \\citep[e.g.][]{2006ApJ...648..652S,2010ApJ...708.1728M}. However, it is also possible that the observed extent reflects the underlying parent body distribution, as might be argued for \\alp, which has dust detected both near and far from the star. Alternatively, the drop in optical depth beyond 50AU required by the continuous \\alp~model may be a sign that the parent planetesimals lie relatively close to the star and that the more distant dust comprises small grains on eccentric and hyperbolic orbits.  Whether debris disks are typically rings or more extended structures is an open question \\citep[e.g.][]{2006ApJ...637L..57K}, in part because obtaining sufficiently high resolution sub-mm observations, those most sensitive to larger grains, is challenging. This ambiguity has only been overcome in a few nearby systems, where a parent body ring has been resolved at sub-mm wavelengths and is seen to be narrower than the radial extent of small grains \\citep{2004Sci...303.1990K,2012ApJ...749L..27W,2012ApJ...750L..21B,2012A&A...540A.125A}. With the development of facilities such as the Atacama Large Millimeter Array (ALMA) and the Northern Extended Millimeter Array (NOEMA), the detection and resolution of larger parent body populations will become possible and will provide insight into the processes that set where planetesimals form and reside.  Looking at the issue of alignment from a wider perspective, coplanarity is the expected outcome for debris disks and planetary systems emerging from the protoplanetary disk phase for small to medium binary separations \\citep[e.g.][]{2000MNRAS.317..773B}. However, only a few examples where the outcome can be tested actually exist. Three systems with transiting circumbinary planets, in which the stars are also eclipsing binaries, show that well aligned systems are a possible outcome \\citep{2011Sci...333.1602D,2012Natur.481..475W}. Further, \\citet{2012Natur.481..475W} find that the occurrence rate of aligned circumbinary planets is probably consistent with the rate for circumstellar planets with similar properties, suggesting that alignment is the typical outcome. However, they acknowledge significant biases exist, and that further work is needed to understand the implications of these discoveries for circumbinary planetary system alignment and frequency.  In the case of circumbinary debris disks, only four systems where the disk and binary alignment can be tested exist; \\alp, \\bet, 99 Herculis \\citep{2012MNRAS.421.2264K}, and HD 98800\\footnote{A weak accretion signature has recently been detected for $\\sim$10Myr old HD 98800 \\citep{2012ApJ...744..121Y}, suggesting that it lies somewhere between the protoplanetary and debris disk phases.} \\citep{2005ApJ...635..442B,2010ApJ...710..462A}. Of these, \\alp~and \\bet~are close binaries with periods of several weeks and HD 98800 has a period of 314 days, while 99 Her has a semi-major axis of 16.5AU and a 56 year period. With only these systems we cannot yet be sure of what trends will emerge. As hinted by the disk-binary alignment of \\alp, \\bet, and HD 98800, and the misalignment for 99 Her, it may be that more widely separated systems are more likely to be misaligned. It will also be interesting to test whether stirring by secular perturbations from binaries is important for disk evolution. This hypothesis could be tested by comparing disk sizes (and structure where possible) with the radii at which secular perturbations can have reached within the system lifetime for a larger sample."
    },
    "1208/1208.3208_arXiv.txt": {
        "abstract": "{We construct an explicit example of a de Sitter vacuum in type IIB string theory that realizes the proposal of K\\\"ahler uplifting. As the large volume limit in this method depends on the rank of the largest condensing gauge group we carry out a scan of gauge group ranks over the Kreuzer-Skarke set of toric Calabi-Yau threefolds. We find large numbers of models with the largest gauge group factor easily exceeding a rank of one hundred. We construct a global model with K\\\"ahler uplifting on a two-parameter model on $\\mathbb{CP}^4_{11169}$, by an explicit analysis from both the type IIB and F-theory point of view. The explicitness of the construction lies in the realization of a D7 brane configuration, gauge flux and RR and NS flux choices, such that all known consistency conditions are met and the geometric moduli are stabilized in a metastable de Sitter vacuum with spontaneous GUT scale supersymmetry breaking driven by an F-term of the K\\\"ahler moduli.} ",
        "introduction": "String theory is a candidate for a fundamental theory of nature since it has the capacity to describe chiral matter fermions with non-Abelian gauge interactions within a consistent theory of quantum gravity. However, at a more detailed level it turns out to be difficult to clearly identify the Standard Model in an expanding universe as one of the possible backgrounds. Part of the problem arises from the poor conceptual understanding of string theory which still is largely based on perturbative formulations. As a consequence there is an abundance of perturbatively consistent backgrounds -- each with a generically high-dimensional moduli space.  This moduli space is particularly apparent if one views the string backgrounds geometrically, that is as compactifications of a ten-dimensional (10D) space-time on some compact six-dimensional manifold. Imposing 4D ${\\cal N}=1$ space-time supersymmetry for phenomenological reasons singles out a specific class of six-manifold which include  Calabi-Yau (CY) manifolds. These manifolds have a large number of non-trivial deformations (moduli) associated with their volume and shape. In a low energy effective description the moduli correspond to 4D massless scalar fields which are flat directions of the scalar potential. Their stabilization has been a long-standing problem but in recent years significant progress has been made at least for certain classes of compactifications.\\footnote{For recent reviews of flux compactifications and their associated uplifts to dS, and a much more comprehensive bibliography, see e.g.~\\cite{Grana:2005jc,Douglas:2006es,Blumenhagen:2006ci,McAllister:2007bg}.}  The mechanism relevant for this paper uses quantized vacuum expectation values (VEVs) for $p$-form gauge field strengths of type IIB string theory \\cite{Dasgupta:1999ss,Giddings:2001yu}. These fluxes generate a scalar potential for a large fraction of the typically ${\\cal O}(100)$ moduli and potentially stabilizes them in a local minimum. The remaining moduli are then fixed by a combination of non-perturbative effects (such as gaugino condensation of 4D ${\\cal N}=1$ gauge theories living on D-branes)~\\cite{Kachru:2003aw}, a combination of perturbative and non-perturbative effects \\cite{Balasubramanian:2005zx}, or an interplay of perturbative effects and negative curvature of the internal space alone~\\cite{Silverstein:2007ac,Caviezel:2008ik,Haque:2008jz,Caviezel:2008tf,Flauger:2008ad,Polchinski:2009ch,Dong:2010pm,Dong:2011uf}.  The mechanism for moduli stabilization can simultaneously break supersymmetry (SUSY) spontaneously by generating non-vanishing $F$- and/or $D$-terms \\cite{Balasubramanian:2004uy,Burgess:2003ic}. Alternatively, SUSY breaking can be achieved by inserting an additional quasi-explicit source, such as an anti-brane in a warped region \\cite{Kachru:2003aw}. The vacuum energy of such fully stabilized compactifications with SUSY breaking can be both positive and negative, leading to a description of de Sitter (dS) space as metastable vacua of compactified string theory~\\cite{Kachru:2003aw,Burgess:2003ic,Balasubramanian:2004uy,Balasubramanian:2005zx,Intriligator:2006dd,Lebedev:2006qq,Haack:2006cy,Cremades:2007ig,Krippendorf:2009zza,Rummel:2011cd,Cicoli:2012fh,Cicoli:2012vw}. The number of these dS vacua is exponentially large due to large number of topologically distinct fluxes necessary for moduli stabilization in the first place. In type IIB string theory compactified on a warped 6D Calabi-Yau manifold there are typically ${\\cal O}(100)$ complex structure or shape moduli associated to the three-dimensional topologically non-trivial subspaces (three-cycles) of the Calabi-Yau. On each three-cycle a flux can be turned on and thus one has ${\\cal O}(100)$ fluxes to choose for stabilizing all the complex structure moduli. For, say, 10 available flux quanta per three-cycle this yields ${\\cal O}(10^{100})$ isolated potential dS vacua~\\cite{Bousso:2000xa,Feng:2000if,Denef:2004ze}. This exponentially large number of dS vacua with a flat number density distribution of the vacuum energy is often called the `landscape' of string vacua. It is coupled to the populating processes of Coleman-deLuccia tunneling and eternal inflation, which together realize space-time regions filled with each dS vacuum infinitely often. As a consequence  Weinberg's anthropic argument for the smallness of the present-day cosmological constant can be realized in string theory. Hence the landscape of dS vacua gave string theory the ability to accommodate recent data from observational cosmology which demonstrated late-time accelerated expansion of our visible Universe, consistent with an extremely small positive cosmological constant $\\Lambda\\sim 10^{-122}\\,M_{\\rm P}^{4}$.\\footnote{The flatness of the vacuum energy distribution on the landscape has recently been questioned by studies using random matrix techniques in general~\\cite{Marsh:2011aa,Chen:2011ac,Bachlechner:2012at}, and a statistical analysis of combined input parameter distributions in the context of K\\\"ahler uplifting~\\cite{Sumitomo:2012wa,3rdUTQUEST2012JAPAN}.}  For the purpose of this work we will restrict ourselves to type IIB warped Calabi-Yau orientifold compactifications with three-form flux which arise as a specific weak coupling limit (Sen limit) of F-theory on elliptically fibered CY fourfolds~\\cite{Vafa:1996xn,Sen:1997gv}. This will allow us to use the known techniques for constructing CYs, calculate their topological data and derive the 4D effective theory. Supersymmetry breaking and lifting the AdS vacuum to dS typically is the least reliable step. Therefore, our goal is to construct explicit global models in the above type IIB context, which exhibit the dynamics of K\\\"ahler uplifting.\\footnote{In this work, we are not addressing the question of constructing a standard model like sector in combination with moduli stabilization, as it has been recently achieved in~\\cite{Cicoli:2011qg,Cicoli:2012vw}. Here, we simply want to avert these complicating features for the sake of moduli stabilization in a stable de Sitter vacuum.} In such models the interplay of gaugino condensation on D7-branes and the leading ${\\cal O}(\\alpha'^{3})$-correction of the K\\\"ahler potential fix the K\\\"ahler moduli in a SUSY breaking minimum, after three-form flux has supersymmetrically stabilized the complex structure moduli and the axio-dilaton. The vacuum energy of this minimum can be dialed from AdS to dS by adjusting the fluxes appearing in the superpotential. Both SUSY breaking and lifting to dS are driven by an F-term of the K\\\"ahler moduli sector which is induced by the presence of the $\\alpha'$-correction~\\cite{Balasubramanian:2004uy,Westphal:2006tn,Rummel:2011cd}. The dS uplift is therefore spontaneous and arises from the geometric closed string moduli. This is the motivation for trying to construct a fully explicit consistent global model including a choice for the flux. Such a model yields an example for a 4D dS space in string theory which is explicit within the limit of the currently available knowledge.  In this work we  discuss constructions of  K\\\"ahler uplifted dS vacua on compact Calabi-Yau manifolds. We perform this analysis both from perspective of F-theory and its weak coupling limit of type IIB string theory compactified on warped Calabi-Yau orientifolds. F-theory arises from the observation that the type IIB axio-dilaton can vary over the compactification manifold $B_3$. One then interprets the axio-dilaton as the complex structure modulus of an elliptic curve fibered over the threefold $B_3$, realizing a complex four-dimensional manifold that takes the form of an elliptically fibered Calabi-Yau fourfold in order to ensure ${\\cal N}=1$ supersymmetry in the 4D effective theory. The attraction of F-theory is due to the geometrization of the complete non-perturbative super Yang-Mills (SYM) dynamics of stacks of 7-branes wrapping four-cycles in type IIB in terms of resolvable ADE type singularities in the fourfold. Of particular interest here are F-theory compactifications on elliptically fibered fourfolds which admit a global weak coupling limit (Sen limit), where the axio-dilaton goes to weak string coupling and becomes approximately constant everywhere on the threefold base of the fibration. In this limit, the set of 7-branes located at the locus where the fiber degenerates,  can be described entirely in terms of perturbative O7-planes and D7-branes.  The analysis in~\\cite{Rummel:2011cd} looked at the effective dynamics of moduli stabilization via K\\\"ahler uplifting. A class of possible examples consists of a Calabi-Yau threefold with an expression for the volume $\\Vol$ given by the sum of one large four-cycle modulus and a collection of smaller so called blow-up four-cycle moduli. This resembles the structure of a swiss cheese with its overall volume given mainly by just the size of the enclosing cycle and many tiny holes (the blow-up four-cycles). We know in addition that in K\\\"ahler uplifted dS vacua the volume $\\Vol$ of the type IIB Calabi-Yau scales with the rank $N$ of the condensing gauge group as $\\Vol \\propto  N^{3/2}$. Moreover, the scale of K\\\"ahler moduli stabilization and thus the resulting K\\\"ahler moduli masses are suppressed by an additional ${\\cal O}(1/\\Vol)$ compared to the scale of the flux-induced complex structure moduli stabilization. These features lead us to search for condensing gauge groups  with a large rank  which induce a large volume.  We will consider models which can be easily uplifted to F-theory compactified on elliptically fibered CY fourfolds that are hypersurfaces in an ambient toric variety~\\cite{Collinucci:2008zs,Collinucci:2009uh,Blumenhagen:2009up}. In addition, we insist on the existence of a smooth Sen limit, as we use the leading $\\alpha'$ correction which is not understood in F-theory in general. The type IIB Calabi-Yau threefold is then the double cover of the F-theory base manifold $B_3$. It is described by an equation such as $\\xi^2=h$, where the orientifold involution is realized by $\\xi\\mapsto -\\xi$. Here $\\xi$ denotes one of the holomorphic complex projective coordinates of the ambient toric variety  where the CY threefold lives, with the orientifold plane sitting at $\\xi=0$. The data describing the D7-brane stacks and the orientifold planes in F-theory via ADE singularities can be specified in terms of sections of holomorphic line bundles and the corresponding homology classes of the associated divisors. The D7-brane tadpole forces the D7-brane to wrap cycles whose homology classes (denoted by $[D7]$) sum up to $[D7]=8[O7]=8[\\{\\xi=0\\}]$ in the type IIB CY threefold. To obtain a large gauge group one generically has to wrap many branes on certain divisors. This is only possible if the coefficients of the orientifold class $[O7]=\\sum_i c_i [D_i]$ are large (here the $D_i$ denote a complete set of divisor four-cycles which in turn form a base of the 2nd Dolbeault cohomology group~$H^{1,1}$). Increasing the homology class $[O7]$ introduces singularities in the base manifold which have to be resolved. Moreover, the Calabi-Yau threefold hypersurface in the weak coupling limit should be free of singularities as well, and therefore the orientifold planes should not intersect each other. This in general turns out to be a severe constraint when one tries to increase the class of the orientifold by choosing  the weights defining the toric variety appropriately.  Finally, we have to check that all used divisors are rigid such that gaugino condensation does contribute to the superpotential. In this context, the role of gauge flux is crucial: On the one hand, a suitable choice can `rigidify' a divisor by fixing some of its deformation moduli \\cite{Martucci:2006ij,Bianchi:2011qh} (see \\cite{Lust:2005bd,Braun:2008pz,Braun:2011zm} for discussions in the F-theory context). On the other hand, switching on gauge flux can generate additional zero modes in the form of chiral matter (especially at the intersection of branes) which  forbids the contribution of gaugino condensation in the superpotential. Moreover, the presence of fluxes can be required by the necessity of canceling the Freed-Witten anomaly~\\cite{Minasian:1997mm,Freed:1999vc}.  As mentioned before, the dynamics of K\\\"ahler uplifting was demonstrated so far on `swiss cheese' type Calabi-Yau threefolds. The cheese with its one big bounding cycle and its many tiny holes implies a certain form of the volume. Let us assume a CY with a set of divisor four-cycles $D_i$ whose (real) volumes we denote by $\\Vol_i$. Then the volume of the a swiss cheese CY is defined by $\\Vol \\sim \\Vol_1^{3/2} - \\sum_i\\Vol_i^{3/2}$ or as $\\Vol \\sim (\\Vol_1 + \\sum_i\\Vol_i)^{3/2} -\\sum_i \\Vol_i^{3/2}$ for an approximately swiss cheese CY. In this situation one can manufacture a large overall volume by enforcing a large gauge group rank on the corresponding divisor $D_1$ which in turn leads to a large $\\Vol_1$. A potential complication may then arise as a high number of branes on $D_1$ typically enforces singularities on other divisors which might yield the overall volume small even though $\\Vol_1$ is large.  At the end, we combine the K\\\"ahler moduli stabilization with an explicit dilaton and complex structure moduli stabilization via RR and NS fluxes. We discuss a specific hypersurface in the weighted projective space $\\mathbb{CP}^4_{11169}$ as a concrete example where the whole program can be executed. In this case, discrete symmetries of the complex structure moduli space  and a specific choice of the three-form fluxes allow us to fix all the complex structure explicitly along the lines of~\\cite{Giryavets:2003vd}. We check that this choice of flux results in values for $W_0$ which are such that the K\\\"ahler stabilization leads to a metastable dS vacuum. Hence, our model constitutes an example for a dS space in string theory which is explicit within the limits of existing knowledge. The only implicitness that remains is the unknown  complex structure moduli dependence of the 1-loop determinant prefactor of the non-perturbative effect. Recent work has shown~\\cite{Rummel:2011cd} that for large volume the mass scale of the K\\\"ahler moduli separates from the scale of the axio-dilaton and the complex structure moduli by one inverse power of the volume. This justifies replacing the complex structure moduli by their VEVs inside the 1-loop determinants, and allows us to parametrize these prefactors as effective constants. Moreover, we can clearly dial the VEVs of the complex structure moduli by availing ourselves of the exponentially large flux discretuum, which easily accounts for a potential mild tuning of the value of the 1-loop determinants.  This paper is organized as follows. In Section~\\ref{lggr_sec} we study the constraints for having a gauge group with large  rank by discussing Kreuzer-Skarke models~\\cite{Kreuzer:2000xy} and hypersurfaces in toric varieties. For the subclass of threefolds with an elliptic F-theory lift ($\\sim 10^{5}$ models) we scan and extract the distribution of the largest-rank gauge group as a function of the number of K\\\"ahler moduli $h^{1,1}$. Then we choose to consider $\\Pex$ as an explicit example and construct large-rank ADE gauge groups on a choice of two divisors, and analyze the consistency constraints both in the type IIB weak-coupling limit, and from the F-theory perspective in sections~\\ref{IIBperspectiveEx} and~\\ref{FthEx_sec}, respectively. Section~\\ref{CSlargeVolume} reviews the general results for supersymmetric flux stabilization of the complex structure moduli and the axio-dilaton. In section~\\ref{dS_sec} we study the scalar potential that stabilizes the K\\\"ahler moduli. We single out a band in the $g_s$ - $W_0$ plane where one finds de Sitter vacua. Here $W_0$ denotes the VEV of the superpotential which arises from supersymmetric flux stabilization of the complex structure moduli. For the explicit model on $\\Pex$, we show how to fix explicitely all the complex structure moduli, thanks to a particular symmetry of the moduli space and a special choice of three-form fluxes. Finally, we check that this choice of flux results in values for $W_0$ which are such that the K\\\"ahler stabilization leads to a metastable dS vacuum. We conclude and discuss our results in section~\\ref{conclusions_sec}. More details of the toric resolution of the $Sp(k)$-singularity can be found in the appendices. We have kept the steps of our calculations rather explicit for future reference but also since certain aspects of the arguments are often  only implicit in the existent literature. ",
        "conclusions": "We discussed the construction of explicit global models in a type IIB context, which exhibit the dynamics of K\\\"ahler uplifting. In this mechanism the interplay of gaugino condensation on 7-branes and the leading ${\\cal O}(\\alpha'^{3})$-correction fixes the K\\\"ahler moduli in a SUSY breaking minimum, after three-form flux has supersymmetrically stabilized the complex structure moduli and the axio-dilaton. The vacuum energy of this minimum can be dialed from AdS to dS by adjusting the flux-induced Gukov-Vafa-Witten superpotential using the flux discretuum. Both SUSY breaking and lifting to dS are driven by an F-term of the K\\\"ahler moduli sector arising from the presence of the $\\alpha'$-correction of the K\\\"ahler potential. Thus the dS uplift is realized entirely by the geometric closed string moduli. This was the motivation for trying to construct a fully explicit consistent global model including explicit flux choice and complex structure moduli stabilization.  In K\\\"ahler uplifted dS vacua the CY volume scales as $\\Vol \\propto N^{3/2}$ with $N$ the rank of the condensing gauge group living on the 7-brane stack wrapping the large four-cycle. Moreover, the scale of K\\\"ahler moduli stabilization and the resulting K\\\"ahler moduli masses are suppressed  by an additional ${\\cal O}(1/\\Vol)$ compared to the scale of flux-induced complex structure moduli stabilization. Hence, we searched for a large gauge group rank to obtain a large volume.  We considered models which can be easily uplifted to F-theory compactifications on elliptically fibered CY fourfolds (embedded in toric spaces). We focused on models that possess Sen's weak coupling limit since we use the leading $\\alpha'$ correction which is not understood for generic points in the F-theory moduli space. Such models have the characteristic feature that the orientifold plane has to be in a homology class of high degree in order to obtain a singularity of large rank $N$.  We checked the consistency conditions to have a globally defined construction, e.g. that the toric base of the elliptically fibered fourfold should be free of singularities of any kind as well as the Calabi-Yau threefold hypersurface in the weak coupling limit. We found that in general this turns out to be a severe constraint when one tries to increase the class of the orientifold by choosing appropriately the weights defining the toric variety.  We then made sure that the volume was of swiss cheese type $\\Vol \\sim \\Vol_1^{3/2} - \\Vol_i^{3/2}$, or at least approximately swiss cheese, e.g $\\Vol \\sim (\\Vol_1 + \\Vol_i)^{3/2} - \\Vol_i^{3/2}$. Then one can manufacture a large overall volume by making $\\Vol_1$ large by enforcing a large gauge group rank on the corresponding divisor $D_1$. We ensure that other large rank stacks wrapping some divisors $D_{i\\not=1}$ (possibly enforced by imposing large rank on $D_1$) did not destroy the large volume approximation.  Finally, we checked that the number of the neutral and charged zero modes could be put to zero, such that the gaugino condensation contribution to the superpotential is non-zero. To do this, the gauge flux on the brane stack has to be chosen appropriately: On the one hand it must be non-zero to `rigidify' the wrapped divisor (if this is not rigid). On the other hand, it should be possible to tune the flux such that it does not generate additional zero modes in the form of chiral matter, charged under the condensing gauge group.  We studied constraints on large gauge group rank by discussing Kreuzer-Skarke models and hypersurfaces in toric varieties. For the subclass of threefolds with an elliptic F-theory lift ($\\sim 10^{5}$ models) we extracted the distribution of the largest-rank gauge group as a function of the number of K\\\"ahler moduli $h^{1,1}$.  Choosing $\\Pex$, which has $h^{1,1}=2$ and $h^{2,1}=272$, as our explicit example we constructed large-rank singularities on a choice of two divisors, and analyzed the consistency constraints both in the type IIB Sen limit, and from the F-theory perspective. The emerging situation for $\\Pex$ looks summarily as follows: We construct an $Sp(24)$ singularity on the `large' divisor $D_1$, which is rigidified by gauge flux, breaking $Sp(24)$ to $SU(24)$. The presence of the $Sp(24)$ stack forces an $SO(24)$ singularity on the second divisor $D_5$, which already is rigid. The gauge flux can be tuned such that no further zero-modes are generated. In F-theory, the gaugino condensation superpotential is related to the superpotential generated by the M5-instantons wrapping the exceptional divisors $E_i$ in the resolved fourfold. In the considered case, we found that the exceptional divisors resolving the $Sp(24)$ singularity satisfy $\\chi_0(E_i)\\geq 1$, that is the necessary condition in the presence of fluxes such that the wrapped M5-instantons contribute to the superpotential.  From the general results for supersymmetric flux stabilization we know, that at the no-scale level the resulting scalar potential is positive semi-definite, which yields full stability of the complex structure sector and the dilaton once fluxes fix them at an isolated supersymmetric point. As K\\\"ahler moduli stabilization via K\\\"ahler uplifting proceeds by breaking the no-scale structure at sub-leading order in the volume (like LVS), the stability of the flux-stabilized complex structure sector extends to the full model. We then analyzed the scalar potential that stabilizes the K\\\"ahler moduli. This singled out a band in $g_s$ - $W_0$ plane where one finds de Sitter vacua. Here $W_0$ denotes the VEV of the Gukov-Vafa-Witten superpotential arising from supersymmetric flux stabilization of the complex structure moduli. The overall volume of the Calabi-Yau threefold was  determined by the data of the construction to be $\\Vol \\sim 52$.  The complex structure moduli space of $\\Pex$ possesses a high-order discrete symmetry $\\Gamma$. We only considered three-form fluxes that respect this symmetry. As a consequence, all the $h^{2,1}-2$ non-invariant complex structure moduli are stabilized. The prepotential of the remaining two complex structure moduli is known via mirror symmetry, and this enables us to stabilize all $h^{2,1}$ moduli explicitly. It is this fact that  in the end allowed us to construct a completely stabilized de Sitter vacuum.  Finally, we gave an explicit flux choice which stabilizes the axio-dilaton and the two $\\Gamma$-invariant complex structure moduli at the right VEVs and value for $W_0$ for the K\\\"ahler stabilization of the explicit construction to proceed into a metastable dS vacuum. In summary, we gave a construction of an example for dS space in string theory which we believe to be explicit and complete within the limits of existing knowledge."
    },
    "1208/1208.5797_arXiv.txt": {
        "abstract": "We present \\emph{Spitzer} IRS spectroscopy of \\cotwo\\ ice bending mode spectra at 15.2 \\micron\\ toward 19 young stellar objects with luminosity lower than 1 L$_{\\odot}$ (3 with luminosity lower than 0.1 L$_{\\odot}$). Ice on dust grain surfaces can encode the history of heating because pure \\cotwo\\ ice forms only at elevated temperature, $T > 20$ K, and thus around protostars of higher luminosity. Current internal luminosities of YSOs with $L < 1$ \\lsun\\ do not provide the conditions needed to produce pure \\cotwo\\ ice at radii where typical envelopes begin. The presence of detectable amounts of pure \\cotwo\\ ice would signify a higher past luminosity. Many of the spectra require a contribution from a pure, crystalline \\cotwo\\ component, traced by the presence of a characteristic band splitting in the 15.2 \\micron\\ bending mode. About half of the sources (9 out of 19) in the low luminosity sample have evidence for pure \\cotwo\\ ice, and six of these have significant double-peaked features, which are very strong evidence of pure \\cotwo\\ ice. The presence of the pure \\cotwo\\ ice component indicates that the dust temperature, and hence luminosity of the central star/accretion disk system, must have been higher in the past. An episodic accretion scenario, in which mixed CO-\\cotwo\\ ice is converted to pure \\cotwo\\ ice during each high luminosity phase, explains the presence of pure \\cotwo\\ ice, the total amount of \\cotwo\\ ice, and the observed residual \\cooo\\ gas. ",
        "introduction": "While the standard star formation model with constant accretion rate \\citep{1977ApJ...214..488S, 1987ARA&A..25...23S} predicts protostars to have a luminosity higher than 1.6 \\lsun\\ most of the time, young stellar objects (YSOs) with luminosity lower than 1.6 \\lsun\\ have long been known \\citep{1990AJ.....99..869K, 1994ApJ...434..614G, 1996ARA&A..34..207H}. The \\emph{Spitzer} Legacy Project, From Molecular Cores to Planet Forming Disks (c2d; \\citealt{2003PASP..115..965E}) found that 59\\% of the 112 embedded protostars have luminosity lower than 1.6 L$_{\\odot}$ \\citep{2009ApJS..181..321E}. In fact, \\emph{Spitzer} found a substantial number of sources with the internal luminosity even below 0.1 \\lsun\\ (\\citealt{2004ApJS..154..396Y, 2008ApJS..179..249D}), which have been called Very Low Luminosity Objects (VeLLOs; \\citealt{2007prpl.conf...17D}). The internal luminosity (\\lint) of an embedded protostar is the luminosity of the central source, excluding luminosity from external heating. Recently, an even lower luminosity protostar ($\\lint < 0.03$ \\lsun) was discovered as a binary companion to IRAM04191+1522 IRS \\citep{2012ApJ...747L..43C}. One explanation for such low luminosities is that mass accretion is not a constant process \\citep{1990AJ.....99..869K, 2009ApJ...692..973E, 2010ApJ...710..470D}. The low luminosity sources may be going through a low mass accretion stage between accretion bursts, explaining their currently low luminosities. If mass accretion is episodic, sources with low luminosity may have had much higher accretion rates in the past, and thus had higher luminosity. Imprints of the high luminosity stage in low luminosity YSOs would support the idea of episodic mass accretion.  Evidence of past periods of higher luminosity can be found in molecular spectra of the gas and ice phases of low luminosity sources (\\citealt{2007JKAS...40...85L, 2012ApJ...754L..18V}). In particular, \\citet{2008ApJ...678.1005P} analyzed the general shape of the 15.2 \\micron~CO$_2$ ice bending mode spectrum toward 50 embedded young stars. The 15.2 \\micron~CO$_2$ spectrum can be decomposed into multiple components, including a pure CO$_2$ ice component. The pure CO$_2$ ice can form by two processes. One is CO$_2$  segregation out of a CO$_2$-H$_2$O mixture. The other is a distillation process, in  which CO evaporates from a CO$_2$-CO mixture, leaving pure CO$_2$ behind. The former process occurs at a high temperature (50 \\-- 80 K) and the latter occurs at a lower temperature (20 \\-- 30 K). Both pure CO$_2$ formation processes are irreversible \\citep{1983A&AS...51..389H}, since the bond between pure CO$_2$ ices makes it the most stable phase. Once the pure CO$_2$ ice has formed, it will not disappear unless it evaporates. Thus the existence of pure \\cotwo\\ ice provides a ``chemical memory\" of warmer conditions in the past.  The total amount of CO$_2$ ice, regardless of form, is another indicator of the temperature history. A number of studies of absorption against background stars have shown that \\cotwo\\ ice must form in the prestellar phase (\\citealt{2005ApJ...627L..33B, 2005ApJ...635L.145K, 2009ApJ...695...94W, 2011ApJ...731....9C}). The shape of the 15.2 \\micron\\ feature toward background stars varies little from cloud to cloud \\citep{2009ApJ...695...94W}, with most ($\\sim 85$ to 90\\%) of the \\cotwo\\ in the \\cotwo-\\water\\ mixture and no evidence for a pure \\cotwo\\ component, despite variations in the ratio of CO to \\water\\ ice from cloud to cloud \\citep{2009ApJ...695...94W}. These ices arise from reactions among H, O, and CO on grain surfaces (e.g., \\citealt{2011ApJ...731....9C}).   Observations of the \\cotwo\\ ice feature toward YSOs tend to show higher values of $N(\\cotwo)$ than the background stars \\citep{2008ApJ...678.1005P}, even for background stars with roughly the same extinction (see Fig. 5 in \\citealt{2011ApJ...730..124C}). The ratio of \\cotwo\\ ice to \\water\\ ice in spectra toward YSOs is also higher and more variable than toward background stars while the ratio of CO to \\water\\ ice is lower for the YSOs, indicative of further processing occuring during the formation of the YSOs \\citep{2011ApJ...730..124C}.  The total amount of CO$_2$ ice formed during the YSO phase depends on the accretion scenario.  The episodic accretion scenario gives multiple long periods of low luminosity \\citep{2010ApJ...710..470D} instead of the very short period of low luminosity that the continuous accretion model predicts \\citep{1977ApJ...214..488S, 2005ApJ...627..293Y}. As a result, episodic accretion provides more time for CO to freeze-out and form \\cotwo\\ ice.  Then, the CO ice can evaporate during episodes of higher accretion and luminosity, leaving the \\cotwo\\ ice behind.  The gas phase CO is the final indicator that we consider. Laboratory experiments show that the binding energies of CO and N$_2$ onto dust grain surfaces are similar \\citep{2005ApJ...621L..33O, 2006A&A...449.1297B}. After including the updated binding energy, chemo-dynamical models \\citep{2004ApJ...617..360L} predict more CO, as measured by rare isotopes like \\cooo, than is observed (e.g., \\citealt{2009ApJ...705.1160C, 2011ApJ...729...84K}). Tying up some of the carbon in \\cotwo\\ ice could help to match the observations \\citep{2012ApJ...754L..18V}.  To explain the observed gas phase molecular lines and the total CO$_2$ ice absorption observed toward one specific low luminosity source, CB130-1-IRS1, \\citet{2011ApJ...729...84K} added some simple ice reactions to an episodic accretion chemical model (\\citealt{2004ApJ...617..360L, 2007JKAS...40...85L}). During low luminosity periods, the gas phase CO freezes into CO ice, and some CO ice turns into CO$_2$ ice, creating a CO-\\cotwo\\ ice mixture; then, during high luminosity phases, the CO evaporates, leaving behind \\cotwo\\ ice, which evaporates only at a still higher temperature ($T \\geq 20$ K), depending on the structure of the ice \\citep{2008ApJ...678.1005P}. This scenario, developed by \\citet{2011ApJ...729...84K} and discussed further by \\citet{2012ApJ...754L..18V}, could explain the current low luminosity, the total amount of CO$_2$ ice, and the strength of the gas phase molecule emission toward CB130-1-IRS1.  In this paper, we test the idea further with higher resolution Infrared Spectrograph (IRS; \\citealt{2004ApJS..154...18H}) spectra for an enlarged source sample and complementary spectroscopy of gas-phase \\cooo\\ at 1.3 mm. We obtained \\emph{Spitzer} IRS, Short-High (SH) mode spectroscopy toward 19 young stellar objects (YSOs) with luminosity lower than 1 L$_{\\odot}$, 4 of them with luminosity lower than 0.1 L$_{\\odot}$. The envelopes around YSOs with luminosity lower than 1 L$_{\\odot}$ have temperatures lower than the CO evaporation temperature. The higher resolution spectra allow decomposition of the 15.2 \\micron~ CO$_2$ ice bending mode spectrum into 5 different components, including a pure \\cotwo\\ component, using the method of \\citet{2008ApJ...678.1005P}. In Section \\ref{sec:sources}, we explain the sample selection in the current study. We describe the observations in Section \\ref{SEC_Observations}. In Section \\ref{sec:analysis}, we describe the analysis of the \\emph{Spitzer} IRS spectrum to get the optical depth and the column density. In Section \\ref{SEC_RESULTS} we present results of the decomposition of ice components. In Section \\ref{sec:Discuss}, we present the setup and results of the chemical models for different accretion scenarios. In Section \\ref{sec:conclusion}, we summarize our findings. ",
        "conclusions": "\\label{sec:conclusion} We detected the pure \\cotwo\\ ice double-peaked feature from 6 out of 19 low luminosity sources  with L $\\leq 1$ \\lsun. Fits to the absorption profile provide strong evidence for pure \\cotwo\\ ice in 3 more sources, for a total of 9 (47\\%). The minimum required temperature to form pure CO$_2$ ice is 20 K, which is not found in any substantial part of the envelope at the current luminosities of these sources.  Detection of pure CO$_2$ ice strongly indicates higher luminosities in the past. During  the time of higher luminosity, the pure CO$_2$ ice forms and persists into the current low luminosity stage. The existence of pure \\cotwo\\ ice toward such low-luminosity sources is very strong evidence of a phase of higher luminosity in the past, consistent with models of episodic accretion. It also presents a challenge to any alternative model for explaining the luminosity spread seen toward YSOs.  In addition to the luminosity spread of sources (\\citealt{2010ApJ...710..470D}, Dunham and Vorobyov 2011), the total CO$_2$ ice amount and the residual gas phase CO are best explained by episodic accretion with conversion of CO to CO$_2$ ice. If 10\\% of CO goes into \\cotwo\\ ice when the gas phase CO freezes, observations of most sources are reasonably well matched. There is a substantial spread of values of both total \\cotwo\\ ice and gas phase \\cooo\\ in the observations, which can be explained in an episodic chemistry. However, variations in the ice formation in the larger cloud can also contribute to this scatter. Determinations of other ice components and extinction in the surrounding cloud can help to distinguish these effects.   This material is based upon work supported by the National Aeronautics and Space Administration under RSA 137730 issued by the Jet Propulsion Laboratory. Jeong-Eun Lee was supported by Basic Science Research Program through the National Research Foundation of Korea(NRF) funded by the Ministry of Education, Science and Technology (No. 2012-0002330). Hyo Jeong Kim and Neal J Evans acknowledge support from NSF grant AST-0607793."
    },
    "1208/1208.0667_arXiv.txt": {
        "abstract": "We present an analysis of the galaxy population of the intermediate X-ray luminosity galaxy cluster, Abell~1691, from SDSS and Galaxy Zoo data to elucidate the relationships between environment and galaxy stellar mass for a variety of observationally important cluster populations that include the Butcher-Oemler blue fraction, the active galactic nucleus (AGN) fraction and other spectroscopic classifications of galaxies. From 342 cluster members, we determine a cluster recession velocity of $21257\\pm54$ kms$^{-1}$ and velocity dispersion of $1009^{+40}_{-36}$ kms$^{-1}$ and show that although the cluster is fed by multiple filaments of galaxies it does not possess significant sub-structure in its core. We identify the AGN population of the cluster from a BPT diagram and show that there is a mild increase in the AGN fraction with radius from the cluster centre that appears mainly driven by high mass galaxies (log(stellar mass)$>10.8$). Although the cluster blue fraction follows the same radial trend, it is caused primarily by lower mass galaxies (log(stellar mass)$<10.8$). Significantly, the galaxies that have undergone recent star-bursts or are presently star-bursting but dust-shrouded (spectroscopic e(a) class galaxies) are also nearly exclusively driven by low mass galaxies.  We therefore suggest that the Butcher-Oemler effect may be a mass-dependant effect. We also examine red and passive spiral galaxies and show that the majority are massive galaxies, much like the rest of the red and spectroscopically passive cluster population. We further demonstrate that the velocity dispersion profiles of low and high mass cluster galaxies are different.  Taken together, we infer that the duty cycle of high and low mass cluster galaxies are markedly different, with a significant departure in star formation and specific star formation rates observed beyond $r_{200}$ and we discuss these findings. ",
        "introduction": "Arguably the most important, emergent conclusion about galaxy evolution during the past decade is that the two main driving forces behind it are the stellar mass of a galaxy and its large-scale environment. Although stellar mass may be the more important of the pair, the environmental effects are also non-negligible and the two variables likely exhibit a significant covariance (Baldry et al.\\ 2006; see also von der Linden et al.\\ 2010).  \\begin{figure*} \\vspace*{-8cm} \\centerline{\\psfig{file=compl2.ps,angle=0,width=6.5in}} \\caption{Completeness for our galaxy sample in terms of apparent r-band magnitude (left) and stellar mass (right). At $r=17.7$ and log(stellar mass)=10.2, the data is $\\sim90$ per cent complete. } \\label{fig:compl} \\end{figure*}   For example, in the extreme (and hostile) environment of the dense cores of massive galaxy clusters galaxy colour is preferentially red and galaxy morphology elliptical -- a systematic difference to non-cluster galaxies (e.g.\\ Dressler 1980; Dressler et al.\\ 1997; Lewis et al.\\ 2002; G{\\'o}mez et al.\\ 2003; Balogh et al.\\ 2004a; 2004b; Kauffmann et al.\\ 2004; Christlein \\& Zabludoff 2005; Baldry et al.\\ 2006; Cooper et al.\\ 2006 Sato \\& Martin 2006; Pimbblet et al.\\ 2006; Bamford et al.\\ 2009; Skibba et al.\\ 2009; Deng et al.\\ 2010; Peng et al.\\ 2010; Ma \\& Ebeling 2011; Patel et al.\\ 2011; Calvi et al.\\ 2012; Jensen \\& Pimbblet 2012).  The mass function of cluster galaxies shows that the more massive galaxies are also redder and more frequently elliptical (e.g.\\ Baldry et al.\\ 2006; di Serego et al.\\ 2006; Clemens et al.\\ 2006; van der Wel 2007; van den Bosch et al.\\ 2008; Bamford et al.\\ 2009; Deng et al.\\ 2010; Pasquali et al.\\ 2010; Peng et al.\\ 2010; Nair \\& Abraham 2010; Thomas et al.\\ 2010; Gr{\\\"u}tzbauch et al.\\ 2011; Giodini et al.\\ 2012; Wilman \\& Erwin 2012; see also Vulcani et al.\\ 2012).  Early investigations in to galaxy evolution inside clusters of galaxies by Butcher \\& Oemler (1978; 1984) suggested that the fraction of blue galaxies decreases rapidly with decreasing redshift. They hypothesized that the cause of this dramatic change in blue fraction to be due to spiral galaxies undergoing ram-pressure stripping caused by interaction with the intra-cluster environment (e.g.\\ Gunn \\& Gott 1972). Such sprials morphologically transform in to lenticular (S0) galaxies and fade as they are no longer able to replenish their gas supplies. Later works confirmed the existence of the Butcher-Oemler effect (Couch et al.\\ 1994; Rakos \\& Schombert 1995; Lubin 1996; Dressler et al.\\ 1997; Smail et al.\\ 1998; Ellingson et al.\\ 2001; Kodama \\& Bower 2001; Margoniner et al.\\ 2001; Goto et al.\\ 2003; Tran et al.\\ 2003; De Propris et al.\\ 2004; Urquhart et al.\\ 2010). Yet observational follow-up by Dressler \\& Gunn (1983) demonstrated that the blue galaxies measured by the Butcher-Oemler effect are actually undergoing star-bursts rather than gently fading away, indicating that instead of gas being stripped in a galaxy, in could be used up in a spectacular, albeit terminal, starburst phase. Couch \\& Sharples (1987) made an independent spectroscopic study",
        "conclusions": "We have presented an investigation of the galaxy population of Abell~1691, an intermediate X-ray luminosity cluster from SDSS data.  Our main conclusions can be stated as follows.\\\\ \\ \\\\ \\noindent $\\bullet$ We have determined new global parameters for Abell~1691, including recession velocity ($\\overline{cz} =21257 \\pm 54$ kms$^{-1}$) and velocity dispersion ($\\sigma_{cz} = 1009^{+40}_{-36}$ kms$^{-1}$).  The cluster is morphologically relaxed in its core regions, with no significant substructure within $r_{200}$ and is `fed' by multiple filaments of galaxies within which a number of sub-groups reside.\\\\ \\ \\\\ \\noindent $\\bullet$ We identify AGN from a BPT diagram and show that the cluster AGN fraction increases with radius from the cluster centre and is mostly driven by massive (log(stellar mass)$>10.8$) galaxies.\\\\ \\ \\\\ \\noindent $\\bullet$ The blue fraction of the cluster population also increases with radius away from the centre but is driven by lower mass galaxies (log(stellar mass)$<10.8$). Emission line galaxies follow a similar pattern. Moreover, the emission line galaxies are found to be driven by low mass galaxies in the cluster (especially the recently star-bursting e(a) class).  This suggests that some low mass galaxies star-burst as they are accreted on to the cluster before fading to redder colours and having their morphologies transformed. Accordingly, we suggest that the Butcher-Oemler effect could therefore be a consequence of mass selection effects.\\\\ \\ \\\\ \\noindent $\\bullet$ Massive galaxies are predominantly red, occupy the centre of the cluster, are non-star-forming, have a different velocity dispersion profile to other cluster galaxies, and are likely to have had a significantly different duty cycle compared to the bluer galaxies.  We have investigated the environment and mass of red and passive spiral galaxies (which are likely to be a transition class object) and show that they are also driven by mass but do not have a preferential position (environment) within the cluster.\\\\ \\ \\\\  This paper follows Jensen \\& Pimbblet (2012) and is our second paper examining the `environment' of intermediate $L_X$ galaxy clusters."
    },
    "1208/1208.0983_arXiv.txt": {
        "abstract": "We analyze line-of-sight atomic hydrogen (\\HI) line profiles of 31 nearby, low-mass galaxies selected from the Very Large Array - ACS Nearby Galaxy Survey Treasury (VLA-ANGST) and The \\HI~Nearby Galaxy Survey (THINGS) to trace regions containing cold (T $\\lesssim$ 1400 K) \\HI~from observations with a uniform linear scale of 200 pc beam$^{-1}$. Our galaxy sample spans four orders of magnitude in total \\HI~mass and nine magnitudes in M$_{B}$. We fit single and multiple component functions to each spectrum to isolate the cold, neutral medium given by a low dispersion ($<$6 km s$^{-1}$) component of the spectrum.  Most \\HI~spectra are adequately fit by a single Gaussian with a dispersion of 8-12 km s$^{-1}$.  Cold \\HI~is found in 23 of 27 ($\\sim$85\\%) galaxies after a reduction of the sample size due to quality control cuts.  The cold \\HI~contributes $\\sim$20\\% of the total line-of-sight flux when found with warm \\HI.  Spectra best fit by a single Gaussian, but dominated by cold \\HI~emission (i.e., have velocity dispersions $<$6 km s$^{-1}$) are found primarily beyond the optical radius of the host galaxy.  The cold \\HI~is typically found in localized regions and is generally not coincident with the very highest surface density peaks of the global \\HI~distribution (which are usually areas of recent star formation). We find a lower limit for the mass fraction of cold-to-total \\HI~gas of only a few percent in each galaxy. ",
        "introduction": "}  Dwarf irregular (dIrr) galaxies in the nearby universe are laboratories for studying the most fundamental properties of gas evolution and star formation.  Large, multi-wavelength studies of these systems are only just beginning.  However, recent, high resolution surveys of large, spiral galaxies in the ultraviolet (e.g., NGS - \\citealt{gild07}), optical (e.g., NFGS - \\citealt{jans00}; ANGST - \\citealt{dalc09}), infrared (e.g., SINGS - \\citealt{kenn03}), and radio (e.g., THINGS - \\citealt{walt08}; HERACLES - \\citealt{lero09}) have given us insight into where and when stars form in a galaxy.  They have also provided clues to the gas conditions surrounding sites of current star formation (\\citealt{lero08}; \\citealt{bigi08}).  The conclusion is that stars form in regions rich in cold, dense, molecular material \\citep{kenn98, kenn12}.  These studies mainly focus on high metallicity systems, however.  How these results translate to low metallicity environments has not been fully explored, mainly due to the fact the most commonly observed tracer of molecular material, CO, is notoriously difficult to observe at low metallicity (e.g., \\citealt{tayl98}; \\citealt{baro00}; \\citealt{lero05}; \\citealt{schr12}).  Nearby dIrr galaxies offer the opportunity to study the star formation process in low metallicity systems at high spatial resolution.  Studies of the star formation rates in dIrr galaxies defined from UV, 24$\\mu$m, and/or H$\\alpha$ emission (e.g., \\citealt{lero08}; \\citealt{bigi08}; \\citealt{royc09}) have found a non-linear correlation with the local atomic hydrogen (\\HI) distributions.  However, the dependence of star formation on molecular gas content is apparent in the large spiral galaxies observed in \\citet{lero08} and \\citet{bigi08}.  The star formation rate surface density is shown to correlate linearly with the molecular gas surface density in these galaxies.  These results can be understood if one assumes that efficient star formation requires molecular hydrogen (e.g., \\citealt{krum09} and references therein).  Based on theoretical studies, the very character of the interstellar medium (ISM) is likely to change at low metallicity (e.g., \\citealt{spaa97}).  \\citet{glov11} modeled the so called X-factor which measures the relationship between amount of detected CO emission and the abundance of molecular hydrogen.  They found that the amount of CO drops substantially at low metallicities, consistent with prior predictions (e.g., \\citealt{malo88}).  This deficit of CO at low metallicities is supported by observations showing relatively low or no detections of CO emission in nearby dIrr galaxies (e.g., \\citealt{tayl98, baro00, lero05, schr12}).  Consequently, this lack of metals also results in lower amounts of polycyclic aromatic hydrocarbon (PAH) emission (e.g., \\citealt{engl05, jack06}). The PAHs are important to the formation of the cold ISM since they are critical in the shielding of UV and soft X-ray photons.  The long standing problem of studying the molecular component in the low metallicity dwarfs continues. We expect dIrr galaxies to contain molecular hydrogen since they {\\it are} forming stars.  However, since direct observations of the molecular gas responsible for forming stars in large samples of dIrr galaxies via CO are currently not feasible, we must find a different tracer of star-forming gas.  One intriguing idea is to find the immediate precursors of the molecular gas.  The assembly of star forming molecular clouds is generally believed to require cold \\HI~(e.g., \\citealt{wolf03}; \\citealt{krum09}).  In our galaxy, cold \\HI~clouds have been observed to surround and even intermix with molecular clouds (e.g., \\citealt{krco10}) through studies of \\HI~narrow self absorption (HINSA; \\citealt{li03}).  One promising technique to find cold \\HI~gas was pioneered by \\citet{youn96, youn97} in a sample of nearby, star forming dIrr galaxies. These authors used high angular and spectral resolution \\HI~data of two nearby dIrr galaxies to decompose line-of-sight spectra into narrow and broad Gaussian components.  The narrow-lined gas ($\\sigma$ $\\sim$ 4.5 km s$^{-1}$) was found only in specific regions within the \\HI~disk while the broad-lined gas ($\\sigma$ $\\sim$ 10 km s$^{-1}$) was found along every line-of-sight observed.  The narrow Gaussian component was attributed to cold (T$\\lesssim$1000 K) \\HI~gas while the broad lined gas was assumed to be warm, neutral hydrogen (T$\\gtrsim$5000 K).  \\citet{youn03} studied three more dIrr galaxies, finding that they, too, contained evidence for cold \\HI~gas.  Other authors have used this technique to discover cold \\HI~in a small number of other galaxies. \\citet{debl06} investigated the \\HI~distribution in NGC 6822 and also found regions rich with cold \\HI~gas.  Recent work by \\citet{begu06} found evidence of cold \\HI~gas in six dIrr galaxies from the Faint Irregular Galaxies GMRT Survey (FIGGS; \\citealt{begu08}).  To date, cold \\HI~in emission has been discovered in 10 nearby dIrr galaxies.  Within the limited data available, different measurements of the cold \\HI~do appear to give comparable results.  Two galaxies, DDO 210 and GR8, from the sample of \\citet{youn03} overlapped with that of \\citet{begu06}. \\citet{youn03} used the Very Large Array (VLA) at a linear scale of $\\sim$200 pc beam$^{-1}$ with a velocity resolution of 1.3 km s$^{-1}$ while \\citet{begu06} used the Giant Metrewave Radio Telescope (GMRT) at a linear resolution of $\\sim$300 pc beam$^{-1}$ with a velocity resolution of 1.65 km s$^{-1}$.  Despite the use of different facilities and different spatial/spectral resolutions, both studies find similar results for both the spatial distributions and velocity dispersions of the narrow component.  This agreement between the two studies suggests that the measurement of cold \\HI~is generally robust, and not extremely sensitive to the exact observing parameters.  The technique of identifying a cold neutral medium has also been used in nearby, high metallicity spiral galaxies.  \\citet{brau97} isolated cold \\HI~gas using relatively low spectral resolution (6 km s$^{-1}$) imaging from the VLA for 11 of the closest spiral galaxies. He combined spectra from discrete radial bins and found that these combined spectra all had a narrow Gaussian core (FWHM $\\lesssim$ 6 km s$^{-1}$) superposed onto broad (FWHM $\\sim$ 30 km s$^{-1}$) Lorentzian wings.  He also found that the cold \\HI~gas is filamentary and is found in clumps, preferentially in the spiral arms.  Since the majority of star formation in spiral galaxies occurs in molecular clouds within the spiral arms (e.g., \\citealt{gord04}), it is not surprising to find the bulk of the cold \\HI~associated with these features.  The purpose of our study is to build upon the previous results of \\citet{youn96,youn97}, \\citet{youn03}, \\citet{begu06}, and \\citet{debl06} in order to provide limits to the locations and amount of cold \\HI~gas in a large sample of 31 nearby, low-mass galaxies using a common spatial resolution.  We describe our galaxy sample in \\S\\ref{obs} and our methods of signal extraction in \\S\\ref{idents}.  Our results are described in \\S\\ref{results} and we end with a summary of our conclusions in \\S\\ref{conclusions}. ",
        "conclusions": "}  We have observed line-of-sight \\HI~emission spectra in 31 nearby, low metallicity dIrr galaxies in order to search for cold \\HI~defined by a velocity dispersion of less than 6 km s$^{-1}$.  We have detected it in 23 of 27 galaxies after quality control cuts were applied. The cold \\HI~may be the future sites of molecular cloud and star formation and to date, are the only way to potentially trace star forming gas at low metallicity.  Based upon our observations, we find:  \\begin{itemize} \\renewcommand{\\labelitemi}{$\\bullet$}  \\item The cold \\HI~gas is found in localized regions and usually at total \\HI~column densities above the canonical threshold of star formation of 10$^{21}$ cm$^{-2}$ \\citep{skil87}. The cold \\HI~in a given galaxy is also not typically associated with the very highest surface density gas (see Figures \\ref{2plots} \\& \\ref{cdplot}).  \\item The cold \\HI~has a typical velocity dispersion of $\\sim$4.5 km s$^{-1}$ (T $\\lesssim$ 800 K).  \\item We derive an average value to the volume filling fraction of the cold \\HI~of 9\\% (assuming our lower limit detections).  \\item We find lower limits to the cold \\HI~gas mass fractions of a few percent, consistent with some models of the multi-phase ISM \\citep{mcke77}.  \\item The SFE$_{HI}$ is roughly constant as a function of the total cold \\HI~mass over the observed gas mass ranges.  \\item The cold \\HI~contributes $\\sim$20\\% of the line-of-sight flux in locations where we detect both a cold and warm component.  \\item Cold \\HI~gas that lacks a warm component is typically found at radii larger than the 25 mag arcsec$^{-2}$ optical radius.   \\end{itemize}  Future work will investigate the relationship between the cold \\HI, local ISM conditions, and tracers of current star formation."
    },
    "1208/1208.0517_arXiv.txt": {
        "abstract": "Using direct numerical simulations (DNS) and mean-field simulations (MFS), the effects of non-uniformity of the magnetic field on the suppression of the turbulent pressure is investigated. This suppression of turbulent pressure can lead to an instability which, in turn, makes the mean magnetic field even more non-uniform. This large-scale instability is caused by a resulting negative contribution of small-scale turbulence to the effective (mean-field) magnetic pressure. We show that enhanced mean current density increases the suppression of the turbulent pressure. The instability leads to magnetic flux concentrations in which the normalized effective mean-field pressure is reduced to a certain value at all depths within a structure. ",
        "introduction": "The Sun's magnetic field is generally believed to be due to a turbulent dynamo operating in the convection zone in its outer 30\\% by radius \\citep{KR80,Mof78,Par79,ZRS83,BS05}. Recent simulations performed by a number of groups have shown that the magnetic field is produced in the bulk of the convection zone. According to the flux tube scenario, most of the toroidal magnetic field resides at the bottom of the convection zone, or possibly just beneath it \\citep{GD00,PM07}. Another possibility is that most of the toroidal field resides in the bulk of the convection zone, but that its spatio-temporal properties are strongly affected by the near-surface dynamics \\cite{KMB12}, or the near-surface shear layer \\citep{B05}. In any case, the question then emerges how one can explain the formation of active regions out of which sunspots develop during the lifetime of an active region.  In the past, this question was conveniently bypassed by referring to the possible presence of a strong toroidal flux belt at the bottom of the convection zone, where they would be in a stable state, except that every now and then they would become unstable, for example to the clamshell or tipping instabilities \\citep{Cally}. However, if the magnetic field is continuously being destroyed and regenerated by the turbulence in the convection zone proper, the mechanism for producing active regions and eventually sunspots must be one that is able to operate within a turbulent environment. One such mechanism may be the negative effective magnetic pressure instability (NEMPI) which is based on the suppression of turbulent pressure by a weak mean magnetic field, leading therefore to a negative {\\it effective} (or mean-field) magnetic pressure and, under suitable conditions, to an instability \\citep{KRR89,KRR90,KR94,KMR96,RK07}. This has been the subject of intensive research in recent years \\citep{BKR10,BKKR12,KBKR12,KBKMR12a,KBKMR12b,KBKMR12,Los12}, following the first detection of such an instability in direct numerical simulations \\citep[DNS; see][]{BKKMR}. Another mechanism that has been discussed in connection with the production of magnetic flux concentrations is related to the suppression of the turbulent convective heat flux \\citep{KM00}. Meanwhile, simulations of realistic solar convection have demonstrated that large-scale magnetic flux inhomogeneities can develop when horizontal magnetic flux is injected at the bottom of the simulation domain \\citep{SN12}, but it remains to be seen whether this is connected with any of the two aforementioned mechanisms.  The purpose of the present paper is to investigate the possibility that higher-order contributions (involving higher spatial derivatives of the mean magnetic field) might play a role in NEMPI. We do this by using DNS to measure the resulting turbulent transport coefficients in cases where a measurable mean current density develops in the DNS. Note that even for an initially uniform mean magnetic field, a mean current density develops as a consequence of NEMPI itself, which redistributes an initially uniform magnetic field into a structured one. To investigate this process further, we use appropriately tailored mean-field simulations (MFS) that show how the spatial variations of the negative effective magnetic pressure vary in space as the instability runs further into saturation. We begin by discussing first the basic equations and turn then to the results. Throughout this work we use an isothermal equation of state which yields the simplest possible system to investigate this process. ",
        "conclusions": "The present results have shown that NEMPI tends to develop sharp structures in the course of its nonlinear evolution. This becomes particularly clear from the MFS presented in \\Sec{MFS}. The results of \\Sec{DNS} suggest that this might have consequences of an intensification of NEMPI with increasing $|\\meanJJ|$, as was demonstrated using DNS. At present it is not clear what is the appropriate parameterization of this effect. One possibility is that the $\\meanJJ$ dependence enters in the same way as the $\\meanBB$ dependence, i.e., $\\overline\\Pi_{xx}=-\\half(1+\\meanJ^2/k_J^2\\Beq^2)\\qp(\\beta)\\,\\meanB^2$, where we treat $k_J$ as a free parameter, although this might be a naive expectation given the small number of data points and experiments performed.  The present results are just a first attempt in going beyond the simple representation of the turbulent stress in terms if the mean field along. Other important terms include combinations with gravity as well as anisotropies of the form $\\meanJ_i\\meanJ_j$. Furthermore, if there is helicity, one could construct contributions to the stress tensor using products of the pseudo-tensors $\\meanJ_i\\meanB_j$ and $\\meanJ_j\\meanB_i$ with the kinetic or magnetic helicity. Such a construction obeys the fact that the Reynolds and Maxwell tensors are proper tensors. Such pseudo-tensors might play a role in the solar dynamo where the $\\alpha$ effect is believed to play an important role. However, nothing is known about the importance or the sign of such effects. It would thus be desirable to have an accurate method that allows one to determine the relevant turbulent transport coefficients."
    },
    "1208/1208.5104_arXiv.txt": {
        "abstract": "We present a pattern-recognition based approach to the problem of removal of polarized fringes from spectro-polarimetric data. We demonstrate that 2D Principal Component Analysis can be trained on a given spectro-polarimetric map in order to identify and isolate fringe structures from the spectra. This allows us in principle to reconstruct the data without the fringe component, providing an effective and clean solution to the problem. The results presented in this paper point in the direction of revising the way that science and calibration data should be planned for a typical spectro-polarimetric observing run. ",
        "introduction": "\\label{sec:intro}  The observation and interpretation of wavelength dependent polarization signals in spectral lines is the primary method for the diagnostics of anisotropic processes in astrophysical plasmas, such as those induced by the presence of deterministic electric and magnetic fields \\cite[e.g.,][]{St94,LL04,CL08,TB10}, or by plasma collisions with collimated beams of ions \\cite[e.g.,][]{Fu08}. At the same time, the depolarizing effects by isotropic collisions and by quasi-random electro-magnetic fields can yield information on the density of the plasma constituents, as well as important insights on the overall complexity of turbulent plasmas at diverse spatial and temporal scales \\cite[e.g.,][]{Ca09b}.  The amplitudes of polarization signals observed in astrophysical plasmas vary widely, ranging from the very small signatures ($\\lesssim 10^{-3}\\,I$, where $I$ is the radiation intensity) typical of the weak-field modifications of scattering polarization by the Hanle effect \\cite[see, e.g.,][]{LL04}, to large amplitudes ($\\gtrsim 10^{-1}\\,I$) induced by the Zeeman effect in the presence of strong magnetic fields, such as those found in sunspots. The weaker polarization signals are easily swamped by systematic errors associated with instrumental effects, which are often difficult to model to the level of precision (\\emph{polarization accuracy}) that is needed in order to isolate the true signals coming from the observed physical system. The calibration of this \\emph{instrumentally induced polarization} is a difficult art. At the same time, the pursuit of ever finer spatial and temporal scales in the investigation of astrophysical plasmas puts continually growing demands on both sensitivity (i.e., signal-to-noise) and accuracy of polarimetric observations \\cite[e.g.][]{Ri08}. Even higher demands are being made by the scientifically critical need to measure chromospheric magnetic fields \\citep{Ju10}.  \\begin{figure}[t!] \\centering \\includegraphics[height=2.9in]{D103517_Sc22_Stk1_true.eps} \\caption{\\label{fig:example} Example of spectro-polarimetric data, showing real spectral line polarization signatures superimposed to instrumental effects, including polarized fringes and detector noise. The abscissa is wavelength, while the ordinate is position along the instrument's slit. The spectral range includes the line of \\ion{Si}{1} at 1082.7\\,nm (around $X=50$), the two components of \\ion{He}{1} at 1083\\,nm (around $X=100$ and $X=130$), the telluric H$_2$O line (around $X=180$), and the \\ion{Ca}{1} doublet at 1083.3\\,nm (approximately from $X=210$ and $X=230$).} \\end{figure}  Polarized fringes are commonly found in spectro-polarimetric data. These are interference patterns that arise because of the presence of optical components (also including air) in a spectro-polarimeter, which have different refractive and/or birefringent properties. Such components include polarization modulators, polarizing beam-splitters, and any optical system where parallel optical interfaces may occur (e.g., interference filters, detector windows). These fringes have the appearance of more or less regular bidimensional patterns, often curved, which typically unfold along the spectral dimension of the data (see Fig.~\\ref{fig:example}). We refer to review studies of polarized fringes \\citep{Se03,Cl04} for a thorough description of this phenomenon. Here, we are interested exclusively with the treatment of this artifact during data reduction.  The treatment of polarized fringes has been a recurring problem for the reduction and analysis of spectro-polarimetric data. The state of the art is to attempt removal of these (and other) instrumental effects using Fourier methods or wavelet analysis \\cite[e.g.,][]{RH06}. Fourier filtering has been successful at removing various types of data artifacts, when their range of spectral and/or spatial frequencies is clearly separated from that of the actual signal of the observed source (L.~Kleint, private communication). Wavelet analysis attempts to identify the dominant frequencies and phases of the fringe pattern in a data frame, when the artifacts are not strictly periodic. \\cite{RH06} have developed a localized solution to fringe pattern reconstruction by employing two-dimensional wavelet transforms, which excel at tracking smooth variations in phase and amplitude of a periodic signal. The shape of a fringe pattern can be isolated in the wavelet space of individually transformed image rows, each row corresponding to a spatial point in the map. An inverse wavelet transform of the fringe space then can reconstruct the fringe pattern in the spatial domain. This can be a particularly powerful method for removing fringes in flat field images, but its application to object images (especially polarized spectral images) is complicated by the contribution of the targeted signal to the local wavelet transform.   Unfortunately, fringe patterns are seldom regular, having an intrinsic bidimensional structure, with variations of the amplitude, frequency, and phase, which can be significant in both the spectral and spatial dimensions (see Fig.~\\ref{fig:example}). Often fringe patterns result from the combination of more than one component, and this combination can also vary smoothly across the dataset. In this case, a two-dimensional wavelet analysis of the observations is challenged by the need to treat every frame separately. In this paper, we propose instead that the identification and removal of polarized fringes might be better approached as a problem of pattern recognition.  Since polarized fringes typically arise within the spectro-polarimeter,\\footnote{A notable exception is represented by fringing that is caused by the polarization calibration optics, which may reside outside the spectro-polarimeter, often well upstream in the optical system of the telescope.} we expect their structure to be predominantly a function of the instrument configuration. Let us consider a scanning slit instrument. During the spatial scan of a given target, the instrument configuration is approximately fixed. We can expect that the polarized fringes will constitute an approximately time-independent pattern within that particular dataset, although there may be some dependence of the fringes' appearance (in both amplitude and phase) on the polarization state of the light entering the spectro-polarimeter. As the spatial scan is acquired, the polarimetric signal of the target will then change over an approximately constant fringe background. Heuristically, an ``orthogonality'' exists between the true polarimetric signal that we wish to analyze and this ``fixed'' fringe background, as a consequence of the fact that the two sources of the line signals and of the background are largely uncorrelated. This suggests that the problem of identification and removal of polarized fringes should be approached as a problem of pattern recognition and feature extraction from a two-dimensional dataset. For this paper we decided to approach this problem using two-dimensional Principal Component Analysis \\cite[2D PCA;][]{Ya04}. Other methods could potentially be adopted instead, which also have been used for the separation of signals from uncorrelated sources, such as the Independent Component Analysis \\cite[ICA; e.g.,][]{JH91}. We defer the study of some of these alternative methods for the problem of the identification and removal of polarized fringes to future work.  The simple fact that spectral line signals and polarized fringe background are largely uncorrelated is central to the success of the PCA approach to the problem at hand. This concept is nicely summarized by \\cite{Jo02} in the Introduction to his book: ``\\emph{The central idea of principal component analysis (PCA) is to reduce the dimensionality of a data set consisting of a large number of interrelated variables, while retaining as much as possible of the variation present in the data set. This is achieved by transforming to a new set of variables, the principal components (PCs), which are uncorrelated, and which are ordered so that the first \\textbf{few} retain most of the variation present in \\textbf{all} of the original variables.}''  Preliminary considerations on the problematics associated with PCA filtering of polarized fringes from Stokes profiles were already advanced in a study on \\emph{compressed sensing} of experimental data \\citep{AL10}. In this work, we provide a novel and more in-depth investigation of the problem. In Sect.~\\ref{sec:PCA} we summarize the main ideas behind two-dimensional PCA, and show how these can be applied to the specific problem of the removal of polarized fringes from spectro-polarimetric data. In Sect.~\\ref{sec:data}, we describe the observations from which the datasets used for the testing of the method were extracted. In Sect.~\\ref{sec:results} we present various examples of application of the method to the datasets previously described, and comment on the quality and reproducibility of the results. Finally, in our concluding remarks, we briefly discuss possible observation and data calibration strategies that could help fully realize the potential of the proposed method. ",
        "conclusions": "The problem of the identication and removal of instrumental artifacts, such as polarized fringes, from spectro-polarimetric maps lends itself naturally to a treatment by pattern recognition methods, e.g., Principal Component Analysis (PCA). In this paper, we have presented the 2D PCA algorithm of \\cite{Ya04}, which appears to outperform other PCA strategies in the tests of the authors of that work.  However, it is interesting to comment briefly also on the more traditional implementation of PCA to face recognition \\citep{TP91}. In this alternate approach, each $m\\times n$ image matrix $\\tens{A}_i$ is rearranged into a $(m\\times n)$-vector $\\bm{B}_i$. The dataset of $N$ images can then be represented by the $(m\\times n)\\times N$ matrix $\\tens{B}=\\{\\bm{B}_i\\}_{i=1,\\ldots,N}$, and the PCA covariance matrix is then built as \\begin{equation} \\label{eq:alternate} \\tens{C}=\\frac{1}{N} \\sum_{i=1}^N (\\bm{B}_i-\\bar{\\bm{B}})^T (\\bm{B}_i-\\bar{\\bm{B}})\\;, \\qquad \\bar{\\bm{B}}=\\frac{1}{N} \\sum_{i=1}^N \\bm{B}_i\\;. \\end{equation} We note that $\\tens{C}$ so defined is an $N\\times N$ matrix, and that both the spatial and spectral dimensions of the original images $\\tens{A}_i$ have been contracted in order to compute it. This approach has the notable advantage of producing eigenfeatures that resemble the original images \\cite[see][for details]{TP91}, which provide a direct visual aid for the selection of the orders that isolate the spectral data from the fringe background. On the other hand, the high level of data compression tends to produce a much slower convergence of the PCA reconstruction series than for the 2D PCA algorithm. The alternative of using the products $(\\bm{B}_i-\\bar{\\bm{B}}) (\\bm{B}_i-\\bar{\\bm{B}})^T$ for the definition of $\\tens{C}$ (see also Sect.~\\ref{sec:PCA}) would preserve both spatial and spectral information of the covariance of the original dataset. On the other hand, this would often create a matrix that is too big to be diagonalized efficiently for the typical values of the image dimensions, $m$ and $n$. That is why the 2D PCA algorithm appears more suitable for our problem, despite the fact that the order selection may be more cumbersome with this approach. We defer a more detailed study of the potential and limitations of the traditional PCA approach to future work.  The strategy to make PCA succeed in the removal of polarized fringes from spectro-polarimetric data is ideally to guarantee the presence in a dataset of widely diverse realizations of a line's polarization profiles over a practically time-independent fringe background. In fact, this determines the condition of non-correlation between the spectral signal and the polarized fringes, which is at the very basis of the working concept of PCA \\citep{Jo02}.  For the \\ion{He}{1} 1083\\,nm data presented in this work, this condition is often met, and as expected PCA manages to separate rather well the spectral line information from the fringe background. The \\ion{Si}{1} line is always prominent in the solar intensity spectrum. However, this should not always be the case for polarization, e.g., for quiet-Sun observations. So one can hope that the same strategy will work with that line as well, at the condition that a sufficient diversity of polarization signals over the fringe background can be acquired during an observation. In the active-region data that we have analyzed, instead, the polarization signals of the \\ion{Si}{1} line are always typically at the level of 10-20\\% for both linear and circular polarization, and thus the subtraction of the fringe background in that spectral range fails consistently. This problem is greatly mitigated by the fact that the typical amplitude of the fringes is relatively small ($\\sim 0.2\\%$) in our data, compared to the observed amplitudes of the \\ion{Si}{1} polarization, and so the polarization profiles of the line are not significantly affected by the fringe background.   The results presented in this paper point in the direction of revising the way that the acquisition of science and calibration data should be planned for a typical spectro-polarimetric observing run. Looking at different targets on the quiet Sun, while maintaining the same configuration of the spectrograph in order to preserve the stationarity of the fringe background, could turn out to be a fundamental addition to all observing programs. Lamp flat fields may also be an important tool for ``fringe calibration'', as they could be used to augment the set of fringe data de-voided of spectral line signals for the purpose of ``training'' the PCA in the identification of fringes. Of course, these lamp flats should ideally be taken under the same optical configuration of the spectrograph as used for a given observation, which may not always be possible."
    },
    "1208/1208.4514_arXiv.txt": {
        "abstract": "X-ray disk winds are detected in spectrally soft, disk--dominated phases of stellar-mass black hole outbursts.  In contrast, compact, steady, relativistic jets are detected in spectrally hard states that are dominated by non-thermal X-ray emission.  Although these distinctive outflows appear to be almost mutually exclusive, it is possible that a disk wind persists in hard states but cannot be detected via X-ray absorption lines owing to very high ionization. Here, we present an analysis of a deep, 60~ksec {\\it Chandra}/HETGS observation of the black hole candidate H~1743$-$322 in the low/hard state.  The spectrum shows no evidence of a disk wind, with tight limits, and within the range of ionizing flux levels that were measured in prior {\\it Chandra} observations wherein a wind was clearly detected.  In H 1743$-$322, at least, disk winds are actually diminished in the low/hard state, and disk winds and jets are likely state-dependent and anti-correlated.  These results suggest that although the launching radii of winds and jets may differ by orders of magnitude, they may both be tied to a fundamental property of the inner accretion flow, such as the mass accretion rate and/or the magnetic field topology of the disk.  We discuss these results in the context of disk winds and jets in other stellar-mass black holes, and possible launching mechanisms for black hole outflows. ",
        "introduction": "Recent {\\it Chandra}/HETGS observations of stellar-mass black holes and neutron stars have revealed disk winds through blue-shifted X-ray absorption lines (Miller et al.\\ 2006a,b; Kubota et al.\\ 2007; Miller et al.\\ 2008; Ueda, Yamaoka, \\& Remillard 2009; Neilsen \\& Lee 2009, King et al.\\ 2012; Ponti et al.\\ 2012, King et al.\\ 2012b).  Such winds are not a negligible part of the accretion flow: estimates for the mass outflow rate in winds range from a fraction of the accretion rate through the disk, to many times greater than the mass accretion rate ($\\dot{m}$) through the disk.  A full understanding of disk accretion now requires an understanding of such outflows, and how they are driven.  More broadly, these winds may provide an important grounding for tentative evidence of ionized winds from the inner disk of AGN (Tombesi et al.\\ 2012; King, Miller, \\& Raymond 2012, King et al.\\ 2012b).  Stellar-mass black hole disk winds appear to be state--dependent: they are detected in spectrally--soft, disk--dominated states, but are not clearly detected in the ``low/hard'' state (Miller et al.\\ 2006a, 2008; Neilsen \\& Lee 2009, Blum et al.\\ 2010; Ponti et al.\\ 2012, King et al.\\ 2012).  In contrast, relativistic radio jets are ubiquitous in the ``low/hard'' state (Fender, Belloni, \\& Gallo 2004), but quenched in disk--dominated soft states (e.g. Russell et al.\\ 2011).  Thus, it is possible that winds and jets are anti--correlated, but related.  An anti-correlation might offer rare clues to the mechanisms that drive wind and jets.  However, the nature of changes to the accretion inflow geometry and radiative processes across state transitions remains unclear (e.g. Esin, McClintock, \\& Narayan 1997; Reis, Fabian, \\& Miller 2010).  Even the apparent absence of disk winds in the ``low/hard'' state may only be an observational effect driven by over-ionization of the gas.  A wind might continue unabated, but simply be impossible to detect through absorption lines owing to a higher ionizing photon flux level.  Ponti et al. (2012) suggested that a higher ionizing flux could not explain the lack of observable wind features in the hard state. However, they did not consider whether the associated change in spectral shape could provide the required change in ionization. Nielsen and Homan (2012) carried out a more complete analysis to show that the extremely dense wind seen in the soft state of GRO J1655-40 by Miller et al. (2006b) could not also be present in a harder state (likely an ``intermediate'' state) observed a few days earlier, with the difference in absorption lines explained by ionization alone. That extreme wind state has only been reported in one other BHB observation, and the driving mechanism may be different from that of the more normal winds in which only Fe XXVI and perhaps Fe XXV are detectable.  In this paper, we compare the Fe absorption lines seen in the soft state of H1743-322 with the lack of absorption lines in a true low/hard state.  We find that the difference in photoionization rate cannot account for the spectra, but that the wind must be genuinely suppressed.  We use a quantitative approach to the difference in absorption line strengths to reach a better understanding of disk winds based on analysis of a deep Chandra HETGS observation of H 1743-322 in the low/hard state, focusing on the physical parameters implied by the apparent absence of the wind.      \\begin{figure} \\includegraphics[scale=0.35,angle=-90]{f1.ps} \\figcaption[t]{\\footnotesize The figure above shows the line spectra from two {\\it Chandra}/HETGS observation of H 1743$-$322.  The data have been divided by a simple continuum and binned for visual clarity.  The observation in black was obtained in a disk--dominated phase; it is listed as ``Observation 1'' in Miller et al.\\ (2006a). A disk wind was detected through the blue-shifted Fe XXV and Fe XXVI absorption lines in that specturm.  The deep low/hard state observation considered in this paper is shown in red.  No significant absorption lines are evident, and restrictive upper limits are obtained through direct fits.} \\end{figure} \\medskip ",
        "conclusions": "A dense, equatorial disk wind was previously detected in {\\it Chandra}/HETGS observations of H~1743$-$322 in disk--dominated states (Miller et al.\\ 2006a).  This source is a strong black hole candidate based on its X-ray properties (e.g. Homan et al.\\ 2005), but its mass and distance have not been determined.  Assuming fiducial values of $d = 8.5$~kpc (based on the proximity of the source to the Galactic center) and $M = 10~M_{\\odot}$, the photoionization models applied in Miller et al. (2006a) describe a wind with $\\dot{M}_{out} =$ 3--4$ \\times 10^{17}~ {\\rm g}~ {\\rm s}^{-1}$ that may originate within $r \\simeq 10^{2-3}~ GM/c^{2}$ of the black hole.  The blue-shifted Fe XXV and Fe XXVI absorption lines found in the brighter, disk--dominated phases are not detected in this deep observation.  Importantly, the ionizing photon flux in this ``low/hard'' state was found to be within the range measured when a disk wind was previously detected.  Using simple arguments, we have shown that the mass outflow rate is likely reduced by at least a factor of $\\sim$3 in the low/hard state, compared to the strongest prior line detections, and might be reduced far more severely.  If a wind with a column density like that measured previously were to persist in the low/hard state, it would would have to be dense and originate close to the black hole, and would likely be driven magnetically (e.g. Miller et al.\\ 2008, Luketic et al.\\ 2010).  A disk wind was absent in H 1743$-$322 during one of the {\\it Chandra} observations in 2003 (see the right-most limits in Figure 2).  Based on photoionization modeling, Miller et al.\\ (2006a) concluded that the non-detection likely required a geometric change (lower density, depth, or covering factor), not merely higher ionization.  Similarly, Neilsen \\& Homan (2012) recently concluded that over-ionization could not account for variability seen in the disk wind in GRO J1655$-$40. Stellar-mass black hole winds may be thermally driven: radiation from the central engine heats gas in the outer disk to the local escape speed (see, e.g., Begelman, McKee, \\& Shields 1983, Woods et al.\\ 1996; also see Luketic et al.\\ 2010).  It is not clear that such winds should be strongly variable.  Geometric changes associated with state transitions --- such as the presence or absence of an {\\it inner} disk (e.g., Esin, McClintock, \\& Narayan; also see Reis, Fabian, \\& Miller 2010) --- need not affect the {\\it outer} disk.  The important parameter for wind detection is the ionizing photon flux, which need not change drastically across states.  Ueda et al.\\ (2010) suggested a geometric change that might explain variability in thermal winds: the development of a hot, geometrically--thick, optically--thick ($\\tau =$7--10) corona that can ``shield'' the outer disk.  This picture may be consistent with the ``very high'' state, and may be able to account for the prior non-detection of a disk wind in H~1743$-$322 (Miller et al.\\ 2006a).  However, this geometrical change is not consistent with the ``low/hard'' state, which is of much greater interest since this is the only state where jets are produced in a steady fashion (e.g. Fender, Belloni, \\& Gallo 2004). Black hole spectra in the ``low/hard'' state require a relatively low optical depth and high temperature when fit with Componization models ($\\tau \\leq$ 1--2, $kT_{e} =$30--120~keV; see e.g. Gierlinski et al.\\ 1997, Torii et al.\\ 2011).  Moreover, the outer disk must be irradiated in order to explain UV and optical emission in the ``low/hard'' state (e.g. Rykoff et al.\\ 2007, Reynolds \\& Miller 2012).  Magnetic driving may provide an alternative to thermal driving, and may be suited to the anti-correlation between winds and jets when comparing ``high/soft'' and ``low/hard'' states.  For instance, the magnetic field configuration may change from toroidal to poloidal in transitions from disk--dominated states to the ``low/hard'' state. Disk winds might then be driven by magnetic pressure generated in a thin disk (e.g. Proga 2003, Ohsuga et al.\\ 2009) while jets might be driven by magneto-centrifugal acceleration along poloidal field lines (Blandford \\& Payne 1982), perhaps aided by black hole spin (Blandford \\& Znajek 1977).  This change could be precipitated by a drop in $\\dot{m}$ through the disk; poloidal fields may be easier to anchor in thicker disks (Reynolds, Garofalo, \\& Begelman 2006).  Alternatively, it is possible that poloidal fields could dominate on each side of a state transition, and that $\\dot{m}_{disk}$ modulates how much mass is loaded onto poloidal field lines (e.g. Spruit 1996).  Winds would originate when the mass outflow rate is high, perhaps breaking field lines or dragging them to make a small angle with respect to the disk (Proga 2003).  Jets would then originate when the mass outflow rate is relatively low, consistent with the ``low/hard'' state, allowing for more effective acceleration.  This latter scenario may be supported by recent work suggesting that black hole winds and jets may be regulated in a common fashion across the mass scale (King et al.\\ 2012b).  The analysis of wind properties across state transtions presented in this work may favor a magnetic wind component.  Stronger support for magnetically--driven winds may derive from photoionization modeling that is able to infer a small launching radius and/or very high mass outflow rate (e.g. Miller et al.\\ 2008), or perhaps from evidence of common regulation of the kinetic power in winds and jets (e.g. King et al.\\ 2012b).  To constrain the launching radius, the density of the gas must be measured directly, since $r^{2} = L / n\\xi$.  Currently, this has only been possible for GRO J1655$-$40, via the detection of the density-sensitive Fe XXII line pair ($n = 10^{14}~ {\\rm cm}^{-3}$), and possibly for NGC 4051 (King, Miller, \\& Raymond 2012).  In contrast, the density of the disk winds in H 1743$-$322 and GRS 1915$+$105 has not been directly constrained (but estimates based on variability are given in Neilsen et al.\\ 2011, 2012).  Trends in these sources may or may not be consistent with thermal driving (Ponti et al.\\ 2012, King et al.\\ 2012b).  Observing stellar-mass black holes with low line-of-sight column densities will facilitate radius and density constraints through the detection of Fe XXII lines, and help to better reveal wind launching mechanisms.   \\vspace{0.2in} We thank the anonymous referee.  We thank Harvey Tananbaum for executing this observation.  ALK acknowledges support through the NASA Earth and Space Sciences Fellowship.  JMM acknowledges support through the {\\it Chandra} Guest Observer program."
    },
    "1208/1208.3272_arXiv.txt": {
        "abstract": "ASTRA (ASTrometric and phase-Referencing Astronomy) is an upgrade to the existing Keck Interferometer which aims at providing new self-phase referencing (high spectral resolution observation of YSOs), dual-field phase referencing (sensitive AGN observations), and astrometric (known exoplanetary systems characterization and galactic center general relativity in strong field regime) capabilities. With the first high spectral resolution mode now offered to the community, this contribution focuses on the progress of the dual field and astrometric modes. ",
        "introduction": "\\label{sec:intro}  ASTRA, the ASTrometric and phase-Referencing Astronomy upgrade funded by the National Science Foundation, aims at expanding the existing capabilities of the Keck interferometer in three incremental steps. First, a self-phase-referencing mode, based on an on-axis fringe tracker makes possible longer integrations for observations with an $R\\sim 2000$ spectral resolution\\cite{Woillez+2010,Pott+2010a,Pott+2010b}, aimed at benefiting YSOs observations\\cite{Eisner+2010a}. Second, a dual-field phase-referencing mode, based on an off-axis fringe tracker will allow observing fainter objects ($K<15$) with a nearby guide star. This mode will mostly benefit extra-galactic astronomy, increasing by an order of magnitude the number of observable AGN. Third, a narrow angle astrometry mode will measure relative positions with precision of $30 \\sim 100$ microarcseconds for objects separated by up to $30$ arcseconds, and further characterize known multi-planet systems. Ultimately, combining the upgraded capabilities of the interferometer with laser guide star adaptive optics on both telescopes will make possible the astrometric monitoring of the inner stars of the galactic center, probing general relativity in the strong field regime. The goals of the project are illustrated in figure \\ref{Fig:MasterFigure}, and the main science cases detailed elsewhere in this proceedings\\cite{Eisner+2010b}.  With the self-phase referencing mode (SPR, section \\ref{Sec:SPR}) in routine operation\\cite{Ragland+2010}, this contribution focuses on the status of the dual-field phase referencing (DFPR, section \\ref{Sec:DFPR}) and astrometric modes (AST, section \\ref{Sec:AST}). For each mode, we will give an overview of the new components added to the interferometer, a summary of which is given in table \\ref{Tab:SubsystemsSummary}, as well as a description of the mode and expected or measured performance. In a last section, we present an astrometric calibration scheme, Differential Narrow angle Astrometry (DNA, section \\ref{Sec:DNA}), that will let us relax the requirements on the knowledge of the astrometric baseline.  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{MasterFigure}\\\\ \\caption{Summary of the ASTRA project and its science cases.} \\label{Fig:MasterFigure} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{SprDfprAst} \\caption{ Architectures of the Self Phase Referencing (SPR, left), Dual Field Phase Referencing (DFPR, center), and Astrometry modes (AST, right). FT-P/S: Primary/Secondary fringe tracker - DLP/S: Primary/Secondary Delay Lines - DFS: Dual Field Subsystem - TTMET: Tip/Tilt Metrology - AMET: Astrometric Metrology - IBM-T: Internal Baseline Monitor Transverse - EBM: External Baseline Monitor - REFB: Reference Baseline } \\label{Fig:ModesArchitectures} \\end{figure}  \\begin{table} \\centering \\begin{tabular}{|c|c|l|c|c|} \\hline {\\bf Mode} & {\\bf Mode status} & {\\bf Subsystem}                               & {\\bf Subsystem status} & {\\bf Section}     \\\\ \\hline \\hline SPR        & operational       & FTS:   Secondary Fringe Tracker Camera        & operational           & \\ref{Sec:SPR}     \\\\ \\hline DFPR       & shared risk       & DFS:   Dual Field Subsystem                   & engineering           & \\ref{SSec:DFS}    \\\\ &                   & TTMET: Tip Tilt Metrology                     & engineering           & \\ref{SSec:TTMET}  \\\\ \\hline AST        & implementation    & AMET:  Astrometric metrology                  & implementation        & \\ref{SSec:AMET}  \\\\ &                   & IBMT: Internal Baseline Monitor - Transverse  & design                & \\ref{SSec:IBMT}  \\\\ &                   & EBM:   External Baseline Monitor              & concept               & \\ref{SSec:EBM}   \\\\ &                   & REFB:  Reference Baseline                     & concept               & \\ref{SSec:REFB}  \\\\ \\hline \\end{tabular} \\vspace{3mm} \\caption{Summary of the ASTRA subsystems, organized by observing mode, with status.} \\label{Tab:SubsystemsSummary} \\end{table}  \\section {SELF-PHASE-REFERENCING}\\label{Sec:SPR}  Now operational, the first step of the ASTRA project was to deliver an on-axis phase referencing mode called self-phase-referencing (SPR) mode. An extensive description of this mode is being published in Woillez et al. (2010)\\cite{Woillez+2010}, along with the first results in Pott et al. (2010)\\cite{Pott+2010a,Pott+2010b} and Eisner et al. (2010)\\cite{Eisner+2010a}. Observations at $R \\sim 2000$ on $K \\sim 7$ targets were demonstrated, as well as the detection of differential phase signature of narrow spectral features with an accuracy of $3 mrad$ on $K \\sim 5$ targets. This mode required the implementation of an additional fringe tracker camera and beam combiner, to allow different integration times between a fast fringe tracking feed stabilizing the optical path difference for a slow spectrograph. From the perspective of the ASTRA instrumentation program, this mode should be considered as a stepping stone to the dual field mode, that exercises the feed forwarding mechanism from the fringe tracker to the science camera and the handling of slow integrations. The architecture of SPR is illustrated on the left side of Figure \\ref{Fig:ModesArchitectures}. For a detailed description of the fringe trackers, developed by the KI project, and extensively used by all the ASTRA modes, see Colavita et al. (2010)\\cite{Colavita+2010}. ",
        "conclusions": "\\end{table}"
    },
    "1208/1208.3758_arXiv.txt": {
        "abstract": "We show that supersonic MHD turbulence yields a star formation rate (SFR) as low as observed in molecular clouds (MCs), for characteristic values of the free-fall time divided by the dynamical time, $t_{\\rm ff}/t_{\\rm dyn}$, the alfv\\'{e}nic Mach number, ${\\cal M}_{\\rm a}$, and the sonic Mach number, ${\\cal M}_{\\rm s}$. Using a very large set of deep adaptive-mesh-refinement simulations, we quantify the dependence of the SFR per free-fall time, $\\epsilon_{\\rm ff}$, on the above parameters. Our main results are: i) $\\epsilon_{\\rm ff}$ decreases exponentially with increasing $t_{\\rm ff}/t_{\\rm dyn}$, but is insensitive to changes in ${\\cal M}_{\\rm s}$, for constant values of $t_{\\rm ff}/t_{\\rm dyn}$ and ${\\cal M}_{\\rm a}$. ii) Decreasing values of ${\\cal M}_{\\rm a}$ (stronger magnetic fields) reduce $\\epsilon_{\\rm ff}$, but only to a point, beyond which $\\epsilon_{\\rm ff}$ increases with a further decrease of ${\\cal M}_{\\rm a}$. iii) For values of ${\\cal M}_{\\rm a}$ characteristic of star-forming regions, $\\epsilon_{\\rm ff}$ varies with ${\\cal M}_{\\rm a}$ by less than a factor of two. We propose a simple star-formation law, based on the empirical fit to the minimum $\\epsilon_{\\rm ff}$, and depending only on $t_{\\rm ff}/t_{\\rm dyn}$: $\\epsilon_{\\rm ff} \\approx \\epsilon_{\\rm wind} \\exp(-1.6 \\,t_{\\rm ff}/t_{\\rm dyn})$. Because it only depends on the mean gas density and rms velocity, this law is straightforward to implement in simulations and analytical models of galaxy formation and evolution. ",
        "introduction": "Stars are the result of the gravitational collapse of cold gas, but on large scales the conversion of cold gas into stars takes much longer than a collapse time. Early explanations for this low star formation rate (SFR) involved the magnetic support against the collapse of giant molecular clouds (GMCs), while the more recent scenario of turbulent fragmentation relies on the support of GMCs by supersonic turbulence \\citep{maclow+klessen04}. However, numerical simulations of star formation in turbulent clouds have generally failed to yield SFRs as low as observed values\\footnote{After submitting this Letter, we became aware of a closely related paper by Federrath and Klessen, submitted to the Astrophysical Journal (arXiv:1209.2856), presenting a large number of star formation simulations that are complementary to ours.} (see \\S~4).  This work shows that supersonic MHD turbulence results in a SFR as low as observed in GMCs, for characteristic values of the free-fall time divided by the dynamical time, $t_{\\rm ff}/t_{\\rm dyn}$, the alfv\\'{e}nic Mach number, ${\\cal M}_{\\rm a}$, and the sonic Mach number, ${\\cal M}_{\\rm s}$ (defined in \\S~2). Using a very large set of adaptive-mesh-refinement (AMR) simulations of driven MHD turbulence, including self-gravity and sink particles, we quantify the dependence of the SFR per free-fall time, $\\epsilon_{\\rm ff}$, on $t_{\\rm ff}/t_{\\rm dyn}$, confirming the results of our previous uniform grid simulations in the case of a very weak mean magnetic field.  The SFR per free-fall time is the efficiency factor of a theoretical Schmidt-Kennicut law, \\begin{equation} \\dot{\\rho}_{\\rm stars}=\\epsilon_{\\rm ff} \\rho_{\\rm gas}/t_{\\rm ff}, \\label{sfr} \\end{equation} often used in models and simulations of star formation in galaxies. We find that $\\epsilon_{\\rm ff}$ decreases exponentially with increasing $t_{\\rm ff}/t_{\\rm dyn}$, has no dependence on ${\\cal M}_{\\rm s}$ alone, and varies by less than a factor of two with variations in ${\\cal M}_{\\rm a}$ within reasonable values for star-forming regions. We therefore propose a new empirical law of star formation that depends only on $t_{\\rm ff}/t_{\\rm dyn}$. ",
        "conclusions": "The new AMR simulations of star formation presented in this work show that i) $\\epsilon_{\\rm ff}$ decreases exponentially with increasing $t_{\\rm ff}/t_{\\rm dyn}$, but is insensitive to changes in ${\\cal M}_{\\rm s}$ (in the range $10\\le{\\cal M}_{\\rm s}\\le20$), for constant values of $t_{\\rm ff}/t_{\\rm dyn}$  and ${\\cal M}_{\\rm a}$. ii) Decreasing values of ${\\cal M}_{\\rm a}$ (increasing magnetic field strength) reduce $\\epsilon_{\\rm ff}$, but only to a point, beyond which $\\epsilon_{\\rm ff}$ increases with a further decrease of ${\\cal M}_{\\rm a}$. iii) For values of ${\\cal M}_{\\rm a}$ characteristic of star-forming regions, $\\epsilon_{\\rm ff}$ varies with ${\\cal M}_{\\rm a}$ by less than a factor of two. We propose a simple law of star formation depending only on $t_{\\rm ff}/t_{\\rm dyn}$, based on the empirical fit to the minimum $\\epsilon_{\\rm ff}$: $\\epsilon_{\\rm ff} \\approx \\epsilon_{\\rm wind} \\exp(-1.6 \\,t_{\\rm ff}/t_{\\rm dyn})$.  This law shows that MHD turbulence is very effective at slowing down the star-formation process and can explain the low average SFR in molecular clouds. Because it only depends on the mean gas density and rms velocity of a star-forming region, the star-formation law we propose is straightforward to implement in simulations and analytical models of galaxy formation and evolution. Future work should test its effect in simulations where the feedback of SNs is accounted for, but the process of star formation is not spatially resolved and needs to be modeled."
    },
    "1208/1208.2274_arXiv.txt": {
        "abstract": "We present the first three-dimensional circulation models for extrasolar gas giant atmospheres with geometrically and energetically consistent treatments of magnetic drag and ohmic dissipation.  Atmospheric resistivities are continuously updated and calculated directly from the flow structure, strongly coupling the magnetic effects with the circulation pattern.  We model the hot Jupiters HD 189733b ($T_{\\mathrm{eq}}\\approx1200$ K) and HD 209458b ($T_{\\mathrm{eq}}\\approx1500$ K) and test planetary magnetic field strengths from 0 to 30 G.  We find that even at $B=3$ G the atmospheric structure and circulation of HD 209458b are strongly influenced by magnetic effects, while the cooler HD 189733b remains largely unaffected, even in the case of $B=30$ G and super-solar metallicities.  Our models of HD~209458b indicate that magnetic effects can substantially slow down atmospheric winds, change circulation and temperature patterns, and alter observable properties.  These models establish that longitudinal and latitudinal hot spot offsets, day-night flux contrasts, and planetary radius inflation are interrelated diagnostics of the magnetic induction process occurring in the atmospheres of hot Jupiters and other similarly forced exoplanets.  Most of the ohmic heating occurs high in the atmosphere and on the day side of the planet, while the heating at depth is strongly dependent on the internal heat flux assumed for the planet, with more heating when the deep atmosphere is hot.  We compare the ohmic power at depth in our models, and estimates of the ohmic dissipation in the bulk interior (from general scaling laws), to evolutionary models that constrain the amount of heating necessary to explain the inflated radius of HD 209458b.  Our results suggest that deep ohmic heating can successfully inflate the radius of HD~209458b for planetary magnetic field strengths of $B\\geq 3-10$ G. ",
        "introduction": "\\label{sec:intro}  Hot Jupiters are unlike any planet in our solar system and their atmospheric circulation exists in an entirely new regime \\citep[for an extensive review, see][]{SCM}.  Orbiting within 0.1 AU of their host stars, these gas giants are subject to stellar irradiation levels $\\sim 10^4$ times the flux Jupiter receives from our Sun.  This leads to a thick radiative zone atop their convective interiors and drives atmospheric winds that in many models reach or exceed the local sound speed \\citep{Showman2009,DobbsDixon2012,Perna2012,RM12b}.  These exotic atmospheres have forced a reevaluation of which of the standard assumptions used in atmospheric dynamics may still be valid and whether new physical processes may need to be included in models in order to accurately determine their circulation patterns.  Although there are a growing number of observed planets, with measurements in expanding wavelength coverage, and an increasingly diverse set of methods by which to characterize their atmospheres, we will still never have as many photons from all exoplanets combined as we do from any single solar system planet \\citep[see][for a review]{Seager2010}.  Nevertheless, these planets provide us with an exciting opportunity to extend the study of planetary atmospheres to strange new worlds.  It has recently been recognized that one of the novel processes that could affect a hot Jupiter's atmospheric circulation is due to the presence of the planet's own magnetic field.  The atmospheres of many of these planets are hot enough that they should be weakly thermally ionized (for pressures near the photosphere), with alkali metals providing the primary source of ions.\\footnote{Note that at lower pressures (from nanobars to microbars) the atmosphere should be ionized by UV radiation from the star.  This is a distinct region from the pressure ranges covered by most general circulation models, with different dominant physical mechanisms and observable signatures.  See \\citet{Koskinen2007} for an example of a circulation model in that regime.}  As the charged particles in the mostly neutral flow are advected around the planet by very fast winds, their circulation through the planetary magnetic field should result in the generation of a secondary component to the field and associated currents.  The effects on the atmosphere are predicted to be a bulk Lorentz force drag on the winds \\citep{Perna2010a} and localized heating from ohmic dissipation of the currents \\citep{Batygin2010,Perna2010b}.  In this paper we present the first circulation models to include these coupled effects.  There are several ways that magnetic effects may influence observable properties of hot Jupiters.  One of the most widely recognized is that the ohmic heating may provide the extra source of heating required to explain the unexpected large radii of some hot Jupiters \\citep{Batygin2011,Laughlin2011,Huang2012,Wu2012}.  For a given planetary magnetic field strength, there should be more heating on planets subject to higher levels of irradiation \\citep{Perna2012}, but these planets will also experience stronger drag on their winds and we may expect an anti-correlation between the amount of radius inflation and the offset of the hot spot from the substellar point \\citep{Menou2012,RM12b}.  There is evidence for temperature inversions in the atmospheres of many hot Jupiters \\citep[e.g.,][and references therein]{Madhu2010} and, while the presence of clouds or stratospheric absorbers will influence the pressure levels at which ohmic heating occurs \\citep{Heng2012}, it is possible that the temperature inversions are in fact produced by ohmic heating \\citep{Menou2012b}.  Finally, with future observing facilities we may be able to directly measure the upper atmosphere wind speeds on these planets and thereby constrain the strength of magnetic drag \\citep{Kempton2012}.  Due to the complexity of this topic, most of the work so far has estimated ohmic heating rates from three-dimensional circulation models or analytic assumptions, without consistently including the feedback of magnetic drag and heating on the circulation pattern.  We have previously published models that attempted to estimate the effect of magnetic drag on hot Jupiter circulation by including a simple, approximate form for the drag \\citep[][also used for Miller-Ricci Kempton \\& Rauscher 2012]{Perna2010a,RM12b}, but here we improve upon that work in several important ways.  First, instead of estimating the drag strength, we now calculate it directly from local conditions (temperature, density) and update it with each timestep, which makes it strongly coupled to the dynamics.  This also means that the magnetic drag is no longer uniform at each pressure level, but instead varies by many orders of magnitude around the planet.  In addition, we include geometric effects due to an assumed aligned dipole field: only the zonal (east-west) component of the flow experiences drag and the strength of drag is also dependent on latitude (with zero drag at the equator).  Finally, we convert the kinetic energy lost through drag into heating from ohmic dissipation.  By including all of these effects, we present the first atmospheric circulation model with geometrically and energetically consistent magnetic drag and ohmic dissipation.  One obvious unknown in studying this topic is the strength of hot Jupiter magnetic fields.  We can use scaling laws to predict field strengths, but our knowledge is limited by the unknown complexity of planet interiors and an incomplete understanding of dynamo theory \\citep[see reviews by][]{Christensen2010,Stevenson2010}.  Nevertheless, such scaling laws estimate hot Jupiter magnetic field strengths to be anywhere from a few to tens of Gauss \\citep[see][and references therein]{Reiners2010}.  Of course it would be preferable to actually measure the magnetic field strength for any particular planet, but unfortunately this will most likely require the use of indirect methods in which the signature of the planet is buried in the stellar signal.  These include the measurement of changes in the stellar chromospheric or X-ray emission due to interaction between its and an orbiting planet's magnetic field (e.g., Shkolnik et al. 2005, 2008, Kashyap et al. 2008; although see Poppenhaeger et al. 2010, 2011, Miller et al. 2012), direct detection of radio emission from the planet \\citep[which is likely too weak for current capabilities, see][and references therein]{Griessmeier2011}, and the signature of a planet's magnetic field in its influence on the rate and geometry of atmospheric evaporation \\citep[e.g.][]{Yelle2004,Adams2011,Trammell2011} or by the field mediating the location of a shock between the planet and its host star's coronal plasma \\citep{Vidotto2010,Vidotto2011}.  We describe our numerical model in Section~\\ref{sec:code}, with particular attention to the new implementation of magnetic effects.  In Section~\\ref{sec:models} we catalog the set of models presented in this paper, explaining our choices for the range of physical parameters used.  We begin the analysis of our results with a detailed look at the model of HD 209458b with $B=3$ G (\\ref{sec:hd2b3}), followed by an examination of models with increasing magnetic field strengths (\\ref{sec:bfield}), and a comparison between our models of HD 209458b and the cooler planet HD 189733b (\\ref{sec:hd1}).  We then discuss the influence of numerical resolution on our results (\\ref{sec:resolution}).  In Section~\\ref{sec:obs} we examine the observable properties of our models and in Section~\\ref{sec:radii} we discuss global ohmic heating rates, commenting upon how they could influence a planet's thermal evolution and its radius.  We conclude with a summary of our main findings in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have updated our three-dimensional model of atmospheric circulation to include geometrically and energetically consistent magnetic drag and ohmic heating.  This is the first model to directly couple the magnetic effects to the full atmospheric structure.  We calculate the electrical resistivities directly from local conditions and update these values at each timestep in the simulation.  We apply magnetic drag only to the zonal (east-west) component of the flow and include the latitudinal dependence in the strength of the drag.  All kinetic energy lost through magnetic drag is consistently returned to the atmosphere as localized ohmic heating.  In these ways the magnetic effects are strongly coupled with the atmospheric circulation.  Here we present results from this code for two well known hot Jupiters, HD 189733b and HD 209458b.  We test models of these planets with planetary magnetic field strengths ranging from 0 to 30 G.  Due to the $\\sim$300 K difference in equilibrium temperature between HD 189733b and HD~290458b, magnetic drag and ohmic heating only influence the circulation of the hotter planet, HD 209458b.  The $B=3$ G model of HD 209458b is obviously different from the non-magnetic version, but even the $B=30$ G model of HD 189733b with enhanced metallicity does not show significant differences from the $B=0$ G model, having only slightly slower wind speeds.  The models of HD 189733b also do not show any signature of magnetic effects in their observable properties, nor do they produce enough ohmic heating to effect the planet's evolution, which is consistent with the observed (non-inflated) radius of this planet.  We find that magnetic effects are able to strongly influence the circulation of HD 209458b; in several aspects the models with $B>0$ differ from the $B=0$ version.  In particular, the magnetic models: 1) have slower wind speeds, by $\\sim$1 \\kms (although still supersonic), 2) do not have an equatorial eastward jet that circles the globe, 3) have departures from hemispheric symmetry in the temperature and flow patterns, and 4) maintain more of a hot-day/cold-night temperature structure, over a deeper range of pressures, than the non-magnetic model.  These trends gradually become stronger at higher magnetic field strengths.  We also find that it is the magnetic drag that has a stronger influence on the circulation than the (coupled) ohmic heating, both of which act primarily on the day side of the planet.  Throughout almost all of the atmosphere the local ohmic heating is a very small fraction of the radiative heating.  However, our $B=30$ G model of HD 209458b has localized features where the ohmic heating reaches as much as 10\\% of the local radiative heating; this model became numerically unstable and crashed before completion.  Magnetic drag and heating have a strong enough effect on the circulation of HD 209458b that they alter its observable properties.  The flux contrast between the day and night sides of the planet is greater in the magnetic models; during the planet's orbit the minimum flux emitted is only $12-13$\\% of the maximum flux, compared to a ratio of 17\\% for the non-magnetic model.  We also find that the brightest region of the atmosphere remains well aligned with the substellar longitude, compared to a 12\\degrees~eastward shift in the $B=0$ G model.  Although there is no shift in longitude in the magnetic models, we do find that the latitude of the brightest region varies in time and can shift away from the substellar point by as much as $10-20$\\degrees.  The differences between the magnetic and non-magnetic models of HD~209458b are at a level such that they could be measured, meaning that models used to interpret observations of this planet should include magnetic effects.  Finally, we compare the ohmic heating profiles from our models of HD 209458b to predictions from evolutionary models to determine whether we find sufficient heating at depth to explain the inflated radius of this planet.  Most of the ohmic heating in our models occurs high in the atmosphere and cannot prevent the planet's standard cooling and contraction.  However, the heating in the deepest layers of our $B=3$ and 10 G models could fulfill the requirement for inflation set by \\citet{Batygin2010}, although only if the internal heat flux is high enough for a hot deep atmosphere; we found heating rates at depth to be two orders of magnitude higher in the $B=3$~G model with $\\Tint=500$ K than the one with $\\Tint=100$ K.  The models with $\\Tint=500$ K also have heating profiles that are increasing with pressure from 10 to 100 bar, even though the winds are becoming slower, with speeds on the order of 100 \\ms.  We use scaling arguments from \\citet{Wu2012} to estimate the ohmic heating in the planet's interior, below our models' bottom boundary.  These estimates meet or exceed the requirements from \\citet{Guillot2002} and \\citet{Batygin2010} for ohmic heating in the adiabatic interior.  Both these estimates, and the rising ohmic power with pressure that we find in our models, suggest that interior ohmic dissipation can inflate the radius of HD 209458b for planetary magnetic field strengths at or greater than $B=3-10$ G.  In order to model the complex interaction between magnetic effects and the atmospheric circulation, we made several simplifying assumptions.  We expect these choices to be appropriate to first order, but as always, we would benefit from a fuller theory, in which these assumptions could be relaxed.  In particular, our calculation of the magnetic drag and ohmic heating is based on a formalism that assumes axisymmetry in the flow structure and atmospheric resistivities \\citep{Liu2008}, neither of which is realized in hot Jupiter atmospheres, especially at low pressures.  We have also assumed the simplest magnetic geometry, with an aligned dipole field, while in reality the planet's magnetic axis could be misaligned from its axis of rotation, or the field could be multipole or uneven in other ways.  Varying the magnetic geometry could have strange impacts on the circulation.  In addition, our magnetic drag is only applied to the zonal flow, although the meridional flow may also begin to experience drag as it nears the poles and the field becomes more radial, but we lack the formalism to include this effect.  However, even with these simplifications, we have established that the inclusion of magnetic effects results in a greater richness of possibilities for atmospheric dynamics on the hottest exoplanets."
    },
    "1208/1208.0115_arXiv.txt": {
        "abstract": "If the lightest supersymmetric particle (LSP) is Higgsino-like, the thermal relic density is lower than the observed dark matter content for a LSP mass in the sub-TeV region. We outline constraints arising from the Fermi Gamma-Ray Telescope data and LSP production from gravitino decay that must be satisfied by a successful non-thermal Higgsino scenario. We show that in a generic class of models where anomaly and modulus mediated contributions to supersymmetry breaking are of comparable size, Higgsino arises as the only viable sub-TeV dark matter candidate if gravitinos are heavy enough to decay before the onset of big bang nucleosynthesis (BBN). The correct relic density can be obtained via modulus decay in these models. As an explicit example, we consider a modulus sector in effective field theory ($D=4,~N=1$ supergravitiy arising from type IIB KKLT compactification). Within this class of mirage mediation models, heaviness of the gravitino forces a sub-TeV Higgsino LSP and gives a Higgs mass around $125$ GeV. In this example, the constraints from direct detection experiments are also satisfied. ",
        "introduction": "Supersymmetry not only stabilizes the Higgs mass against quantum corrections, it also provides a candidate for dark matter. In $R$-parity conserving models the lightest supersymmetric particle (LSP) is stable, hence a dark matter candidate. The lightest neutralino, which is a mixture of Bino, Wino and Higgsinos, is the most suitable dark matter candidate with the prospect for detection in various direct and indirect searches.  In this work we point out that a comprehensive solution to the cosmological gravitino problem motivates the dark matter to be Higgsino-like. Gravitinos heavier than ${\\cal O}(40)$ TeV have a lifetime shorter than 0.1 s and decay before the onset of big bang nucleosynthesis (BBN). This results in a considerable relaxation as the gravitino abundance will not be subject to tight BBN bounds~\\cite{BBN}.  In effective supergravity, the masses of the Bino and Wino are sensitive to the mass of the gravitino $m_{3/2}$  \\cite{Kaplunovsky:1993rd}, and in particular, for $m_{3/2} > 40$ TeV, one typically has Bino and Wino masses above TeV in type IIB modulus mediation models. On the other hand, the Higgsino mass depends on the $\\mu$ parameter, which can be reduced by anomaly mediated contribution to supersymmetry breaking. As a result, if we demand that the dark matter particle has a mass in the sub-TeV region, the Higgsino becomes a more natural candidate.   If the lightest neutralino is pre-dominantly Higgsino, with mass in the sub-TeV region, the annihilation rate is typically larger than the nominal value $\\langle \\sigma_{\\rm ann} v \\rangle = 3 \\times 10^{-26}$ ${\\rm cm^3 s^{-1}}$, thus resulting in an insufficient thermal relic abundance~\\cite{Baer:2012uy}. A natural way to obtain the correct dark matter relic density is to consider non-thermal sources of Higgsino production.  We consider scenarios where Higgsino dark matter is non-thermally produced by a late decaying modulus \\cite{Moroi:1999zb}. We find that for annihilation rate to be compatible with bounds from the Fermi Gamma-Ray Telescope~\\cite{fermi}, the modulus decay should reheat the universe to a temperature $T_{\\rm d} \\sim {\\cal O}({\\rm GeV})$. An additional requirement is that the branching ratio for modulus decay to the gravitino is $\\ls {\\cal O}(10^{-5})$, so that the decay of gravitinos thus produced does not lead to dark matter overproduction.  As an example of the modulus sector, we consider the standard scenario of KKLT compactification \\cite{Kachru:2003aw}, with the Kahler modulus reheating the universe around $1$ GeV. Within this framework, for appropriate values of the relative contributions of anomaly and modulus mediated contributions, Higgsino emerges as the dark matter candidate. The annihilation rate is consistent with the Fermi bounds, and the correct relic density is obtained by non-thermal production. The Higgs mass $m_h \\sim 125$ GeV \\footnote{Recent experiments at the LHC have provided strong hints of a Higgs-like particle at $\\sim 125$ GeV \\cite{LHCHiggs}.} is also satisfied in this scenario, and we find that it actually requires the gravitino mass to be in the cosmologically safe region.  Moreover, the spin independent scattering cross section is consistent with the latest bounds from direct detection experiments~\\cite{direct}.  Within this specific example, however, decay of the gravitinos that are directly produced from modulus decay overproduces dark matter. This is a direct consequence of the couplings between the modulus and the helicity $\\pm 1/2$ components of the gravitino, which are in turn set by the underlying K\\\"ahler geometry of the effective $D=4,~N=1$ supergravity theory. We summarize a set of geometric conditions in the modulus sector that are sufficient to ensure consistent non-thermal Higgsino dark matter as outlined above.  We note that apart from the purely cosmological motivations shown in this study, the Higgsino also emerges as the LSP within the framework of Natural Supersymmetry as discussed in \\cite{Papucci:2011wy}, \\cite{Hall:2011aa}, \\cite{Baer:2012uy}.   The paper is organized as follows. In Section \\ref{Gravproblem}, we relate the cosmological gravitino problem with the preference for Higgsino dark matter. In Section \\ref{nonthermal}, we outline the conditions that must be satisfied by any successful scenario of non-thermal Higgsino dark matter. In Section IV, we work out an explicit example of a non-thermal scenario. In Section \\ref{gravitino}, we outline the general constraints on an effective modulus sector in order to avoid overproduction of gravitinos. We conclude the paper in Section VI. ",
        "conclusions": "Considering dark matter in the sub-TeV range, thermal freeze-out underproduces Higgsino-like LSP. It is possible to enhance the relic density using non-thermal mechanisms of dark matter production. The enhancement, however, needs to obey the constraints from the gamma ray flux from dwarf spheroidal galaxies and the dark matter content of the universe. Moreover, there should be no overproduction of dark matter at any later stage, for example by the decay of gravitinos.  In this paper, we demonstrated these ideas in a generic class of mdoels where anomaly and modulus mediated contributions to supersymmetry breaking are comparable. Interestingly, we found that within this class of models, heavy gravitinos that are not subject to BBN bounds force the Higgsino as the only viable dark matter candidate in the sub-TeV range. We considered an explicit example of mirage mediation model in $D=4,~N=1$ supergravitiy arising from type IIB KKLT compactification, where the modulus decay provides the non-thermal origin of Higgsino-like dark matter. The large gravitino mass is helpful to yield $m_h$ around $125$ GeV in this model and satisfy the constraints arising from the dark matter direct detection experiments. We also discussed the general conditions to avoid the overproduction of LSP from gravitino decay in such scenarios."
    },
    "1208/1208.6426.txt": {
        "abstract": "We study the effect that uncertainties in the nuclear spin-dependent structure functions have in the determination of the dark matter (DM) parameters in a direct detection experiment. We show that different nuclear models that describe the spin-dependent structure function of specific target nuclei can lead to variations in the reconstructed values of the DM mass and scattering cross-section. We propose a parametrization of the spin structure functions that allows us to treat these uncertainties as variations of three parameters, with a central value and deviation that depend on the specific nucleus. The method is illustrated for germanium and xenon detectors with an exposure of 300 kg yr, assuming a hypothetical detection of DM and studying a series of benchmark points for the DM properties. We find that the effect of these uncertainties can be similar in amplitude to that of astrophysical uncertainties, especially in those cases where the spin-dependent contribution to the elastic scattering cross-section is sizable. ",
        "introduction": "\\label{sec:introduction}  \\let\\oldthefootnote\\thefootnote \\renewcommand{\\thefootnote}{\\alph{footnote}} \\footnotetext[1]{MultiDark Fellow} \\footnotetext[2]{MultiDark Scholar} \\let\\thefootnote\\oldthefootnote \\setcounter{footnote}{0}  Direct searches of dark matter (DM) aim to observe this abundant but elusive component of the Universe by detecting its recoils off target nuclei of a detector (for a recent review, see, e.g. Ref.\\,\\cite{Cerdeno:2010jj}). A large number of experiments have been taking data in the last decades or are currently under construction with this objective, leading to a very exciting present situation.   In fact, some experiments have claimed potential signals that could be compatible with the detection of a weakly-interacting massive particle (WIMP). This is the case of the DAMA collaboration \\cite{Bernabei:2003za}, which observed an annual modulation in the recoil rate on a NaI target that was later confirmed by the upgraded DAMA/LIBRA detector \\cite{Bernabei:2008yi}. Similarly, the CoGeNT collaboration, with a germanium target, reported an irreducible excess in their data that could point towards very light WIMPs \\cite{Aalseth:2010vx} and also observed an annual modulation effect \\cite{Aalseth:2011wp} although the latter is not easy to reconcile with the DAMA/LIBRA result. Finally, the CRESST-II experiment, which uses CaWO$_4$ as a target, also reported an excess \\cite{Angloher:2011uu} over the expected background. However, these observations are in conflict with the negative results obtained in searches by other experimental collaborations. Experiments such as CDMS-II \\cite{Ahmed:2009zw,Ahmed:2010wy}, XENON10 \\cite{Angle:2011th}, XENON100 \\cite{Aprile:2011hi,Aprile:2012}, EDELWEISS \\cite{Armengaud:2011cy}, SIMPLE \\cite{Felizardo:2011uw}, KIMS \\cite{Kim:2012rz}, and a combination of CDMS and EDELWEISS data \\cite{Ahmed:2011gh} are in strong tension with the regions of the parameter space compatible with WIMP signals in DAMA/LIBRA or CoGeNT. Moreover, a reanalysis of CDMS data has been performed in order to look for annual modulation with negative results \\cite{Ahmed:2012vq}.   The elastic scattering cross-section of WIMPs off nuclei can be separated in two components, spin-independent (SI) $\\sigsinu$, and spin-dependent (SD) $\\sigsdnu$, which originate from different terms in the Lagrangian describing the interaction of a DM particle with quarks. The SI term stems from scalar or vector couplings and its contribution to the total WIMP-nucleus cross-section scales as the nucleon number squared, $A^2$, whereas the SD term originates from axial-vector couplings and its total contribution to the cross-section off nuclei is only a function of the total nuclear angular momentum and the DM spin. Thus, the SI term typically dominates for heavy nuclei.   Constraints are normally expressed in terms of the SI and SD components of the WIMP-nucleon elastic cross-section, $\\sigsi$ and $\\sigsd$, respectively. To date, the most stringent constraints on $\\sigsi$ are those obtained from the XENON100 data \\cite{Aprile:2012} that exclude SI cross-sections above $\\sigsi \\approx 2 \\times 10^{-8}$ pb for a mass around 50 GeV, as well as XENON10 \\cite{Angle:2011th} and the low-energy reanalysis of CDMS-II \\cite{Ahmed:2010wy}, which dominate for light WIMPs. Regarding the SD contribution, the leading bounds from direct detection experiments have been provided by XENON \\cite{Angle:2008we} (SD cross-section with neutrons, $\\sigsdn$) and COUPP \\cite{Behnke:2010xt} and PICASSO \\cite{Archambault:2012pm} (SD cross-section with protons, $\\sigsdp$) but indirect detection experiments such as SuperKamiokande \\cite{Tanaka:2011uf} and IceCube \\cite{Abbasi:2009uz}, as well as searches for mono-jet \\cite{Goodman:2010yf,Goodman:2010ku} and mono-photon plus missing energy in Tevatron \\cite{Aaltonen:2012jb} and LHC \\cite{Rajaraman:2011wf,Chatrchyan:2012pa,Fox:2011pm}  lead to even more compelling constraints on $\\sigsdp$. Larger and more sophisticated direct detection experiments are currently under development that will be able to explore the DM parameter space with unprecedented sensitivity. This is the case, for example, of the SuperCDMS and XENON1T collaborations, which aim at the construction of 1 Ton scale detectors based on germanium and xenon, respectively.   In the light of this promising experimental situation, it seems plausible that the DM can be discovered in the near future in direct detection experiments. In such an event, the study of the signal rate and spectrum (differential rate) can be used to determine some of the DM properties, namely its mass, $\\mwimp$, and elastic scattering cross-section \\cite{Green:2007rb,Green:2008rd,Drees:2007hr}. The precision of this reconstruction is very sensitive to the characteristics of the detector and is affected by uncertainties in the parameters describing the DM halo, as well as in the nuclear form factors. Astrophysical uncertainties have been widely discussed in the literature \\cite{Green:2010gw,Akrami:2010dn,Green:2011bv,Strege:2012kv} and they are known to introduce significant errors in the determination of the mass and scattering cross-section of DM. Regarding nuclear uncertainties, those in the SI form factor have been argued to be relatively small \\cite{Chen:2011im}. The effect of variations in the SD form factors has not been previously addressed and constitutes the objective of this work.   We consider the hypothetical future observation of a DM candidate in a direct detection experiment and, sampling over the three-dimensional space of $(\\mwimp,\\,\\sigsi,\\,\\sigsd)$, we investigate how the reconstruction of these quantities is affected by nuclear uncertainties in the spin-dependent structure function of the target nucleus. In order to do so, we propose a description of structure functions based on three parameters, which enlarge the parameter space sampled, and allow us to incorporate uncertainties in a consistent and systematic way. This provides a general method, applicable to any detector target. We particularize our analysis for the case of a germanium detector (such as, e.g., SuperCDMS), for which we consider the spin-dependent structure functions provided by the analysis of various groups \\cite{Bednyakov:2006ux,Ressell:1993qm,Dimitrov:1994gc}, and for xenon detectors (such as, e.g., the future XENON1T), for which we use the structure functions derived in Ref.\\,\\cite{Ressell:1997kx} and \\cite{Menendez:2012tm}.   We observe that the effect of nuclear uncertainties in SD structure functions can lead to variations in the reconstructed DM mass and SD elastic cross-section, the effect being more important in those scenarios in which the SD term in the WIMP-nucleus cross-section is the main contribution to the total detection rate. In such cases uncertainties in the spin-dependent structure functions are similar in amplitude to those induced by astrophysical uncertainties in the DM halo parameters, although the latter also affect the SI component.   The paper is organized as follows. In Sec.\\,\\ref{sec:DD} we introduce the formalism used to compute the recoil event rate, emphasizing the role of SD interactions. We concentrate on the case of a germanium detector, introduce the models available in the literature that describe the spin-dependent structure function and comment on their differences. Sec.\\,\\ref{sec:Bayesian} describes the generation of the simulated data for a set of benchmark models, and the implementation of the scanning algorithm to probe the phenomenological parameter space. In Sec.\\,\\ref{sec:results} we show the reconstruction of DM parameters for each benchmark scenario, using different nuclear models for the SD structure function and investigating how this alters the predictions for the DM properties. In Sec.\\,\\ref{sec:parametrization} we present a parametrization of the SD structure function that allows us to systematically account for uncertainties when scanning over our parameter space, and we apply the method to the cases of germanium and xenon detectors. Our conclusions are summarised in Sec.\\,\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   We have studied the effect that uncertainties in the nuclear spin-dependent structure functions have in the reconstruction of DM properties by means of direct detection experiments.   Assuming a hypothetical future observation of DM in a direct detection experiment we have systematically investigated how well its phenomenological parameters $(\\mwimp,\\,\\sigsi,\\,\\sigsd)$ can be determined when uncertainties in the SD form factors of the target nuclei are taken into account. We focused at first on the case of a germanium target and considered two possible models describing the SDSF of its isotope $^{73}$Ge, sensitive to SD WIMP couplings. Using a Bayesian inference algorithm we determined for each of these models the pdf and profile likelihood of the DM parameters in a set of benchmark scenarios. We observed that if a model is chosen to describe the SDSF of a particular nucleus, the reconstruction of the DM properties can strongly depend on the choice made (see in this sense the comparison between the predictions using the R- or D-model in Figs.\\,\\ref{fig:BM3_profl} and \\ref{fig:BM3_pdf}). In particular, differences in the reconstructed values of the WIMP mass as well as the SD component of the WIMP-nucleon scattering cross-section appear. In general these effects are more important when the SD contribution to the total detection rate is not negligible.   In the second part of the paper we have proposed a description of the SD structure functions in terms of three parameters which fit the zero-momentum value and the slope of the SDSF, and account for the presence of a high-momentum tail. This allows us to include uncertainties in the SDSF in the sampling of the parameter space and treat them in a consistent and systematic way. Using this method we have computed the profile likelihood for the DM parameters for the same three benchmark points as before, in the case of a germanium-based and a xenon-based detector.   Finally, we have explicitly compared the effect of nuclear uncertainties in the SDSF with those that are associated with the parameters of the halo of dark matter. We find that uncertainties in the SDSF can even be comparable in magnitude to astrophysical ones when the SD contribution to the total detection rate is sizable.   \\vspace*{2ex} \\noindent {\\bf Acknowledgements}\\\\[0.5ex] We have greatly benefited from discussions with A.M.~Green, B.~Kavanagh and L.~Robledo. We are grateful to A.~D\\'iaz-Gil for technical support with the computing facilities at the Instituto de F\\'isica Te\\'orica and H.D.~Kim and the Seoul National University for allowing us to use their computational facilities in the last stages of the project. D.G.C. is supported by the Ram\\'on y Cajal program of the Spanish MICINN. M.F is supported by a Leverhulme Trust grant. J.-H.H. is supported by a MultiDark Fellowship. M.P. is supported by a MultiDark Scholarship. This work was supported by the Consolider-Ingenio 2010 Programme under grant MultiDark CSD2009-00064. We also thank the support of the Spanish MICINN under grant FPA2009-08958, the Community of Madrid under grant HEPHACOS S2009/ESP-1473, and the European Union under the Marie Curie-ITN program PITN-GA-2009-237920.  \\clearpage"
    },
    "1208/1208.2132.txt": {
        "abstract": "Dynamical dark energy (DE) phenomenon emerges as a geometrical effect accompanying the cosmological expansion of nonrelativistic fermionic matter. This occurs without the need for any fluid, like e.g. dynamical scalar field (quintessence, cosmon, etc.), and with conventional form of the Einstein equations   in  contrast to other known geometrical DE models.  The phenomenon results from first principles in the framework of the two measures field theory where,  in the Einstein frame, both fermion masses and the cosmological constant (CC)  turn into functions of the  cold fermion density $n$. This $n$ dependence  becomes negligible  in regular (laboratory) conditions but it may have an important role in cosmology. In the 4D gravity model where the original action involves only CC  and massive fermions without selfinteraction,  for different (but wide) regions in the parameter space we have found two possible classes of scenarios for the late universe starting from the cold matter domination era. We argue that the fermions which drive the variable CC should be  associated with cold neutrinos disposed in voids and supervoids. The cosmological dynamics of the first class practically coincides with that of the $\\Lambda$CDM model, while the dynamics of the second class is of the phantom-like regime with a pseudo-rip scenario. Crossing the phantom divide happens due to a new type of the neutrino DE effect where  neutrinos pass through the state with zero mass and with the vacuum-like EoS $P_{\\nu}=-\\rho_{\\nu}$. ",
        "introduction": "Numerous attempts to understand the nature of the discovered cosmic acceleration\\cite{acceleration} have originated the common feeling in the cosmology and particle physics community that even though all observational data are in agreement with the simplest model where this phenomena is described by the cosmological constant (CC), the latter is not enough to give adequate answers to \\textit{all} challenges of  the  cosmology. Therefore to explain the origin of the acceleration it seems to be impossible to advance without some kind of changes in fundamental of modern physics\\cite{Dolgov}. The new physical entity responsible for the present accelerated expansion of the universe was denominated as  the dark energy (DE) and the problem was exposed  intensive study with  vast of models,  for reviews see\\cite{{Sahni-Starobinski-review}}-\\cite{DEreview_Odintsov}. Using terminology of the review\\cite{{Sahni-Starobinski-review}}, to a large extent, all models with variable DE can be  divided into two main classes: 1) physical DE models where DE is associated with some fluid usually realized by means of a new, very weakly interacting physical field.  2) geometrical DE models where effect of DE is realized by means of modified gravity.  In the most of the physical DE models, like e.g. quintessence\\cite{quint}, coupled quintessence\\cite{amend}, mass varying neutrinos and neutrino DE \\cite{MassVar}-\\cite{GK_neut_DE}, growing neutrinos\\cite{growing-neutr}, chameleon cosmology\\cite{chameleon}, phantom\\cite{phantom}, quintom\\cite{quintom}, the DE entity is modelled by scalar field(s) with very special features (very flat potential, etc.) or/and using  non-scalar systems, e.g. models with specially adapted spinors\\cite{spinorDE} . Using of such fields means the exit from the framework of the matter content of the particle Standard  Model.  The geometrical DE models, e.g.  higher dimensional DGP braneworld DE \\cite{DGP},\\cite{Chimento_Maartens}, $f(R)$ DE \\cite{f(R)}, scalar-tensor models\\cite{scalar-tensor}, braneworld with Einstein-Gauss-Bonnet gravity\\cite{Ein-Gauss_Bonne},   etc. (see also reviews \\cite{Sahni-Starobinski-review},\\cite{Modified Gravity and Cosmology},\\cite{DEreview_Odintsov}) demonstrate possibilities to realize an effective DE without any field or other matter entity intended to produce DE. However the price for this effect is the emergence of nonconventional 4D gravitational dynamics. Only after certain regrouping of terms having a geometrical meaning to the effective matter source terms, the gravity equations take the Einsteinian GR form.  The model of the present paper is realized in the framework of the two measures field theories (TMT)\\cite{GK1}.  We want to stress that  TMT possesses a number of peculiar features which allow to regard it as  a new type of alternative theories. First, the modification of the gravity and matter sectors in TMT is implemented on a common basis. Second,  apart from proceeding in the first order formalism, all the modification in the action of TMT , as a matter of principle,  reduces to the use  of the volume element with two measures: the regular one $\\sqrt{-g}$ and the new one $\\Phi$ involving additional non-dynamical degrees of freedom (for details see Sec.II).  Third, TMT admits but does not require to involve in the underlying Lagrangian terms and degrees of freedom different from those of the conventional theory (i.e. Einstein gravity and Standard particle model). Fourth, in the Einstein (physical) frame, the gravity  and all matter field equations have conventional form without the need in regrouping and renaming geometrical and matter source terms. The novelty consists in emergence of variable cosmological constant, very nontrivial form of scalar field potential, fermion selfinteraction and  Yukawa type coupling constants become local function of the fermion density. Therefore fermions \\textit{generically} may possess very unconventional properties that makes it possible, for example, to realize effect of the neutrino driven DE which is the subject of this paper. But  properties of  \\textit{fermionic matter in regular conditions  are undistinguished from those in the conventional theory}. The term 'regular conditions' means here that the local fermion energy density $\\rho_f$ is tens orders of magnitude larger than the vacuum energy density $\\rho_{vac}$, that is actually a condition which is fulfilled  in all  measurements carried out so far in physics, including tests of GR. For example in regular conditions,  corrections to fermion mass  have the order $\\rho_{vac}/\\rho_f$. In this respect, in contrast to other alternative theories, the novelty of the most of the studied so far TMT models is that in order to achieve for undesirable deviations from the conventional theory to be unobservable there is \\textit{no  need in special tuning or constraints on the masses and coupling constants in the underlying Lagrangian}.  In our earlier publications  on TMT\\cite{G1}-\\cite{{emerging}},\\cite{GK_neut_DE} we concentrated on studying the scale invariant TMT models.  Here the scale invariance is spontaneously broken by the integration of the equations of motion and we have found that the dilaton plays the role of the inflaton field in the early universe and that of the DE field in the late universe. It was shown\\cite{GK4} that the Yukawa type effective coupling constant $f$ of the dilaton to matter depends on the local matter density.  For matter in regular conditions the magnitude of $f$ is suppressed by the factor of the order of $\\rho_{vac}/\\rho_f$. The model yields such  resolution of the fifth force problem automatically, without any special tuning of the parameters.  In Ref.\\cite{GK_neut_DE} devoted to cosmological applications of the model, we have studied the effect of neutrino DE emerging in the course of evolution of the late time universe filled with the homogeneous dilaton field and the cold gas of uniformly distributed non-relativistic neutrinos. A new kind of regime in the cosmological dynamics have been revealed where neutrinos undergo transition to a state with unbounded grows of their mass and this process is accompanied with  reconstruction of the scalar field potential in such a way that this dynamical regime appears to be energetically more preferable than it would be in the case of the universe with no fermions at all.  Recently another TMT model\\cite{G} involving a gravitating scalar field with a Born-Infeld type kinetic term and with an arbitrary potential have been constructed.  This model gives a unified picture of dark energy and dark matter reproducing the $\\Lambda$CDM picture. A further generalization, including the addition of a tachyon\\cite{} allows the formulation of an \"inverse quintessence\" scenario.  In the present paper we work with the TMT model which may be regarded as a simplified TMT modification of the Standard Model. This is the case because even though the underlying Lagrangian of our simplified model involves only free massive fermions and the CC term, the TMT allows\\cite{GK2-1} to extend it by adding all fields and symmetries of Standard Model without changing the main results of the paper. The reason is that the key role in generation of the aforementioned effects of the novelty of TMT (in comparison with conventional theory) belongs to the algebraic constraint which appears due to the new measure of integration $\\Phi$. The constraint determines the scalar field $\\zeta =\\Phi/\\sqrt{-g}$ as a local function of $\\overline{\\Psi}\\Psi$ (where $\\Psi$ refers to massive fermion fields) and the latter is the only contribution to the constraint from the Standard Model fields. Therefore, at least in the first step of studying the cosmological applications of the model, one can restrict ourself with the simplified model.  In spite of the fact that aside from fermions, our simplified model does not contain any additional dynamical degree of freedom, the variable effective CC $\\Lambda_{tot}$ is generated and its local space-time value is determined by the local fermion density. In this paper we explore the cosmological dynamics of the late spatially flat FLRW universe starting from the cold matter domination epoch. $\\Lambda_{tot}$ is governed by the density of the cold fermion gas involved in the cosmological expansion. The main goal of this paper is to demonstrate that in the TMT modification of the Standard Model, a variable CC is  generated without the need for additional dynamical degrees of freedom. To our knowledge, this is the first model where this idea is realized, see for example Refs.\\cite{Vilenkin},\\cite{VCC}.   The organization of the paper is the following. In Sec.II we present: a review of the basic ideas of TMT; discuss properties of the new degrees of freedom; describe  on qualitative level some details of using the action principle resulting in the equations of motion and the constraint in terms of the original set of variables; explain, again qualitatively, the reasons and the results of transition to the Einstein frame. This allows us, after the formulation of the model in Sec.III, to present the equations of motion in the Einstein frame at once. Sec.IV is devoted to the very detailed analysis of physically permitted intervals for the geometrical scalar field $\\zeta$. In Sec.V we describe new aspects of the cosmological averaging procedure caused by the appearance of the constraint. Afterwards we present equations of cosmological dynamics resulting from our field model after cosmological averaging and classify possible scenarios of the late universe. In Sec.VI the results of numerical solutions are collected and classified in accordance with qualitative analysis of Sec.V. In Sec.VII we shortly summarize main results of the model that should help  to the reader to see the picture altogether and finally present some preliminary speculations concerning the possibility  to describe the fermion dark matter as a distinctive fermion state in TMT. ",
        "conclusions": "\\textit{I. Fermions in TMT.}  One of the peculiarities of the TMT consists in the perfectly novel features of massive fermions, and the geometrical scalar field  $\\zeta$, Eq.(\\ref{zeta}), plays the main role in this respect:  1. The fermion mass is $\\zeta$ dependent, while the local value  $\\zeta(x)$ is determined via the constraint as a function of the local fermion density. As the local fermion energy density $\\rho_f$ is many orders of magnitude larger than the vacuum energy density $\\rho_{vac}$, the fermion mass is constant up to corrections of the order of $\\rho_{vac}/\\rho_f$, and the local value  $\\zeta(x)$ must be very close either to $\\zeta_1$ or to  $\\zeta_2$ (see the step II of Sec.IV). Thus, such \"high density\" fermion states belong to the regular particle physics situation.  2. As $\\rho_f$ becomes comparable with $\\rho_{vac}$, the fermion mass becomes $\\zeta$ dependent, and due to the constraint this means that the fermion mass depends upon the fermion density. For short, we called such kind of fermion state as the low density one.  It follows from the last two items that in order to describe the fermion state in the quantum field theory formulated in the framework of TMT, in addition to the usual quantum numbers one should add the value of $\\zeta$.   3. In addition to the canonical energy-momentum tensor, fermions possess also the dynamical fermionic $\\Lambda^{(f)}_{dyn}$ term,  Eq.(\\ref{Tmn-noncan}), and therefore the appropriate noncanonical pressure $P_{(f,noncan)}=\\Lambda^{(f)}_{dyn}$ and the noncanonical energy density $\\rho_{(f,noncan)}$ satisfy  the equation  $P_{(f,noncan)}=-\\rho_{(f,noncan)}$. We have shown that  $\\rho_{(f,noncan)}$ and $\\rho_{vac}$ are \\textit{always} of the same or very close orders of magnitude. For that reason the effect of the noncanonical energy-momentum tensor is unobservable in all particle physics measurements till now but it has a crucial role in the cosmology of the late time universe.   \\textit{II. Dynamical DE imitated by the geometrical effect of TMT in the presence of massive fermions.} The matter content of the 4D TMT model studied in the present paper is reduced to massive fermions though radiation could be added without altering the results because gauge fields do not contribute to the constraint. The important feature of the model is that it does not involve any fluid intended to describe the DE. In the case of the absence of fermions, it follows from the constraint that $\\zeta =const$ and then the model contains the constant CC. However, in the presence of fermions, on account of the constraint, $\\zeta$ becomes a function of the local fermion density. In its turn, the constant CC turns into the  $\\zeta$ dependent CC, i.e. on account of the constraint, the CC becomes dependent of the fermion density. This CC together with $\\Lambda^{(f)}_{dyn}$ form the effective  variable CC, $\\Lambda_{tot}$, Eq.(\\ref{T_DE}) depending upon the fermion density. Note that since  $\\Lambda_{tot}$ contains the term $\\Lambda^{(f)}_{dyn}$ it would be incorrect, as usually, to treat $\\Lambda_{tot}$ as the variable {\\it vacuum} energy density. It seems to be more natural to regard $\\Lambda_{tot}$ as the effective dynamical DE imitated by the geometrical effect of TMT in the presence of massive fermions and driven by the fermion density.  \\textit{III. Cosmological averaging and cosmology of the late universe.}  1. We have concentrated on the cosmological outputs of the key effect of the  model: generating of the variable CC $\\Lambda_{tot}$. We have shown that the constraint keeps its form after the cosmological averaging and in the course of the cosmological expansion, the cosmological average of $\\Lambda_{tot}$ with the very high accuracy may be regarded as a function of the  density $n$ of cold fermions disposed in the maximal volume domains with the low fermion density. We have proposed to associate these fermions with neutrinos because it seems to be natural to assume that, in the course of cosmological expansion, from all known fermions only nonrelativistic neutrinos are able to survive without interactions up to the state when their canonical (thermal) pressure is negligible in comparison with their noncanonical pressure $\\Lambda^{(f)}_{dyn}$ . If this assumption is true then the dynamical DE effect is driven by neutrinos disposed in voids and supervoids.  We have found the  permitted by TMT intervals of $\\zeta$ where cold fermions can evolve continuously from  the high density  to the low fermion density state ended with asymptotic transition $\\rho\\rightarrow \\rho_{vac}$. According to these intervals of $\\zeta$ we have classified possible scenarios of the late universe. Almost in all types of scenarios (except for the I.A type), $P_{(f,noncan)}=\\Lambda^{(f)}_{dyn}$ is negative which means that $\\Lambda^{(f)}_{dyn}$ describes {\\it the neutrino DE} and the latter is the essential fraction of the total effect.  2. For the I.B type of scenario, where the {\\it total} EoS $w=p/\\rho$ decays \\textit{monotonically} from $w\\approx 0$ to $w= -1$, there exists a wide range in the parameter space such that the differences in the behavior of the total EoS and of the total averaged energy density from those  in the $\\Lambda$CDM model are practically unobservable.  3.   In another wide region of the parameter space, the model admits  a phantom-like cosmological scenario where the crossing of the phantom divide is caused by changing the sign  of the effective neutrino mass in the course of the cosmological expansion. As the result, at this moment the cold neutrinos energy density and pressure satisfy exactly the vacuum EoS that means disappearance of neutrinos as particles (excitations) in a sense this term has in the quantum field theory. Such a limiting realization of the idea of the neutrino DE \\cite{Neutr-dark}-\\cite{GK_neut_DE} deserves more detailed investigation. Even though the sign of fermion mass in the relativistic quantum mechanics is a question of choice, it remains unclear the physical meaning of  the new cosmological dynamics effect we observe in our model:  \"degeneration\" of neutrino with positive effective mass with subsequent  \"regeneration\" with negative effective mass. Nevertheless this strange effect, from our viewpoint, seems to be less dangerous than the use of the phantom scalar field\\footnote{ Notice that the total measure $\\Phi +k\\sqrt{-g}$  in the fermion kinetic term of the underlying action (\\ref{totaction}) does not change the sign at the moment of crossing the phantom divide. In this point the model of the present paper is drastically different from the model studied in Refs.(\\cite{GK3}), (\\cite{GK6}). There we have studied the scale invariant TMT model involving dynamical scalar field (dilaton) which plays the role of the inflaton in the early universe and the role of DE in the late time universe. In that model, the crossing of phantom divide occurs at the moment when the cosmological evolution of the dilaton yields the change of the sign of the total measure in the dilaton kinetic term of the underlying action, i.e. the phantom scalar field cosmology emerges in the course of dynamical evolution instead of introducing the wrong sign kinetic term into the action, as usual. }. We should note also that in this paper we restrict ourself with studying massive fermions without taking into account the Higgs mechanism. Preliminary study shows that in TMT,  in the context of cosmology, a possible   back reaction of fermions on the vacuum expectation value of the Higgs field may be relevant. In this respect the  model  of the present paper may be regarded as a toy one. It would be interesting to investigate in the future what kind of effect  such a back reaction may have on the change of the sign of the neutrino mass and on the possibility of the phantom-like scenario.     \\textit{IV. The present value of the vacuum energy density}. The effect of the dynamical DE  was realized here on the basis of the first principles of TMT. The only additional quantitative assumption is that the dimensionless parameters $b$, $k$ and $h$ have the same order of magnitude. As we mentioned in the step I of Sec.IV, the observed tiny value of $\\rho_{vac}$ may be achieved if these parameters are huge numbers. We want to argue here that the latter choice is very much different from what is usually called a fine tuning in resolution of the old or new CC problem\\cite{Weinberg1}. When the CC problem  one hands over to particle field theory, the generally accepted understanding of the required subtraction or cancellation mechanism consists in a fine tuning of {\\it coupling constants and masses in   similar  terms in the Lagrangian}, like e.g. as it is done in SUSY and SUGRA. The way large dimensionless numbers we would like to appear in the TMT action is absolutely different: 1) the parameters $b$, $k$ and $h$ are huge numbers as compared with dimensionless parameters in particle field theories but there is no need in mutual fine tuning of  $b$, $k$ and $h$; 2) huge values of $b$, $k$ and $h$ have no any observable consequences except for the tiny value of the vacuum energy density in  comparison with the typical local matter energy density.   \\textit{V. Generalization to the Standard Model matter content}.  In our earlier papers\\cite{GK2-1} we have shown that the gauge fields may be added to the model without changing the constraint. The nonabelian symmetries are also compatible with TMT. Therefore the matter content and symmetries of the model may be extended up to those of the Standard Model.  We would like to stress again that results of the present paper allow us to claim that the effect of the dynamical DE  may be realized with the only matter content of the Standard Model and without any additional dynamical degrees of freedom intended to get this effect. Note that  all the peculiar TMT effects (new with respect to the conventional theory) emerging in the presence of fermions become unobservable in the regular particle physics conditions.    \\textit{VI. Fermion dark matter candidate?} One can finally propose the hypothesis concerning the nature of the dark matter. Creation of fermions in regular particle physics conditions implies that the fermion emerges in the \"high density\" regime. As we have mentioned in the item I of this section, in such a case the fermion may be in the state with $\\zeta$ very close either to $\\zeta_1$ or to  $\\zeta_2$. In our earlier papers\\cite{GK2-1} concerning studying the Standard Model in the framework of TMT, we explored the idea that fermion states with $\\zeta$ very close to $\\zeta_1$ and  $\\zeta_2$ might be associated with the first and second fermion generations respectively. However in the present paper we run into a drastic difference in features of the $\\zeta_1$ and  $\\zeta_2$ fermion states. In fact, we have elucidated that if fermions have been created in the state with $\\zeta$ very close to $\\zeta_1$ and  satisfying the conditions (\\ref{interval_zeta_A}) or (\\ref{interval_zeta_B}) then they are able to evolve to the arbitrary low energy density. If  however fermions have been created in the state with $\\zeta$ very close to $\\zeta_2$ then it follows from Eq.(\\ref{sign_bigger_zero}) that  only in two following cases it is possible: (1) $\\zeta_0<\\zeta_2<\\zeta<\\zeta_1$ and (2)  $-b<\\zeta<\\zeta_2<\\zeta_0<\\zeta_1$ \\footnote{ The case (2) involves an additional requirement on the parameter $b$ which is satisfied in the numerical solutions of Sec.VI for the scenarios Ib, II and III.}. One can check using Eq.(\\ref{diff-eq-zeta}) that $\\zeta$ is forced to oscillate in the interval $\\zeta_2<\\zeta<\\zeta_1$ in the case (1) and in the interval $-b<\\zeta<\\zeta_2$ in the case (2). Such regimes should perhaps be ended at some intermediate equilibrium values, but for us it is important here that the fermion starting from such a state is unable to evolve to a state with arbitrary low energy density. We come to this conclusion in the FLRW metric but one can expect a similar result e.g  in a spherically symmetric gravity problem.    Since $\\zeta$ is supposed to be a continuous function, one can expect that there should be a mechanism which prevents a possibility  for the fermion in the states of the cases (1) and (2) to be in the same space-time region as well as for each of them to be in the same space-time region with another fermion whose state starts from  $\\zeta$ satisfying the conditions (\\ref{interval_zeta_A}) or (\\ref{interval_zeta_B}). Suppose that the dark matter consists of the same fermions (from the viewpoint of the quantum field theory) as the visible fermion matter. Our hypothesis may be formulated in the following form:   the fermion in the state  with $\\zeta$ satisfying the conditions (\\ref{interval_zeta_A}) or (\\ref{interval_zeta_B}) is the visible fermion while the same fermion but in the state as in  the cases (1) or (2) is a candidate for the dark matter particle. It may be that inability to exists in the same space-time region explains the absence of the visible to dark matter local interactions. Besides, if for example in the case (2), the equilibrium value that $\\zeta$ reaches in the above mentioned oscillations  is close to $-b$ then the  mass of the  fermion  in such a state may be much bigger than the one of the visible fermion state (see Eq.(\\ref{rhofcan})). Then physics of such fermions may be very much different from the regular particle physics. It would be interesting   to investigate a possibility to realize  conditions for an existence of a heavy sterile neutrino."
    },
    "1208/1208.2056_arXiv.txt": {
        "abstract": "Many models of ultra-high energy cosmic-ray production involve acceleration in linear accelerators located in Gamma-Ray Bursts magnetars, or other sources.   These source models require very high accelerating gradients, $10^{13}$~keV/cm, with the minimum gradient set by the length of the source. At gradients above 1.6 keV/cm, muons produced by hadronic interactions undergo significant acceleration before they decay.  This acceleration hardens the neutrino energy spectrum and greatly increases the high-energy neutrino flux.  We rule out many models of linear acceleration, setting strong constraints on plasma wakefield accelerators and on models for sources like Gamma Ray Bursts and magnetars. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.5440_arXiv.txt": {
        "abstract": "We analyze and model an M1.0 flare observed by \\textit{SDO}/AIA and \\textit{Hinode}/EIS to investigate how flare loops are heated and evolve subsequently. The flare is composed of two distinctive loop systems observed in EUV images. The UV 1600 \\AA\\ emission at the feet of these loops exhibits a rapid rise, followed by enhanced emission in different EUV channels observed by AIA and EIS. Such behavior is indicative of impulsive energy deposit and the subsequent response in overlying coronal loops that evolve through different temperatures. Using the method we recently developed, we infer empirical heating functions from the rapid rise of the UV light curves for the two loop systems, respectively, treating them as two big loops of cross-sectional area 5\\arcsec\\ by 5\\arcsec, and compute the plasma evolution in the loops using the EBTEL model \\citep{klim08}. We compute the synthetic EUV light curves, which, with the limitation of the model, reasonably agree with observed light curves obtained in multiple AIA channels and EIS lines: they show the same evolution trend and their magnitudes are comparable by within a factor of two. Furthermore, we also compare the computed mean enthalpy flow velocity with the Doppler shift measurements by EIS during the decay phase of the two loops. Our results suggest that the two different loops with different heating functions as inferred from their footpoint UV emission, combined with their different lengths as measured from imaging observations, give rise to different coronal plasma evolution patterns captured both in the model and observations. ",
        "introduction": "It is widely accepted that magnetic reconnection occurs in solar flares with a large amount of energy released in the corona \\citep{prie02}. By reconnection, magnetic field lines change connectivity and form new loops. The released energy is transported downward along the loop by non-thermal electrons or thermal conduction to produce enhanced optical and ultraviolet (UV) emissions at the footpoints of the loop. In the meanwhile, impulsive energy deposition in the lower atmosphere evaporates hot plasma into the newly formed flare loops. The plasma then cools down, becoming visible sequentially as loops in soft X-ray and extreme-UV (EUV) emissions.  Assuming that energy deposit takes place during the impulsive phase of the flare, some modelers adopt a loop heating rate proportional to the observed hard X-ray light curve and model heating and evolution of the flare as one single loop (e.g., \\citealt{fish90, raft09}). However, it has been recognized that most flares consist of multiple loops formed and heated successively \\citep{hori98, asch05, mulu11}. With this consideration, \\cite{warr06} applied multi-thread 1D hydrodynamic simulations to a flare, and was able to reproduce the soft X-ray radiation of the flare as observed by \\textit{GOES} and \\textit{Yohkoh}. To match observations, the author used 50 threads (flare loops), each heated at a different time for about 200 s by a prescribed heating rate. Similarly, \\cite{hock12} have modeled a two-phase flare observed by the Atmospheric Imaging Assembly (AIA; \\citealt{leme12}) and EUV Variability Experiment (EVE; \\citealt{wood12}) on board the \\textit{Solar Dynamics Observatory} (\\textit{SDO}). They observed different flare loop systems brightened in these two phases, and used the 0D model called ``enthalpy-based thermal evolution of loops\" (EBTEL; \\citealt{klim08}) to compute plasma evolution in these loops with arbitrary heating rates that are adjusted to match with the observed EUV radiation. In their study, 22 loops are heated throughout the flare to match the observed flare evolution.  Key to all loop heating models is the heating function, which describes where, when, for how long, and by how much a flare loop is heated. To date, most multi-thread loop heating models prescribe an arbitrary number of loops, each given an arbitrary amount of heating at an arbitrary time, and the parameters are adjusted to best match the observed light curves of coronal emission. To relate flare heating with magnetic energy release in a more realistic manner, \\cite{long10} developed a model to compute the rate of energy released in reconnection using the time-dependent reconnection flux measured from observations of flare ribbon evolution. The amount of released energy is then used to compute flare plasma evolution and coronal radiation and compare with observations by \\textit{GOES} and \\textit{RHESSI}. Recently, \\cite{qiuj12} introduced an alternative method to observationally infer heating rates by analyzing spatially resolved UV light curves emitted at the feet of over a thousand flare loops anchored to brightened UV footpoints. The loop heating occurs when the UV emission starts to rise, the duration of heating is twice the rise time of the UV light curve, and the magnitude of the heating rate is assumed to be directly proportional to the UV counts measured at the foot of the loop. With this method, many properties of the heating rate are observationally constrained.  For its high efficiency in modeling a large number of loops, a number of studies have used the EBTEL model to compute plasma evolution in flare loops, as listed in Table \\ref{comtab}. These studies by different groups treat the model inputs and outputs in different ways. The input to the model includes the number of flare loops and their properties (such as the loop length) and their heating functions. The model-calculated temperature and density of flare loops are then used to compute X-ray and EUV emissions by flare loops that can be compared with observations. Furthermore, the EBTEL model includes a mean enthalpy flux between the corona and its base at the transition region, which describes upflows predominant during the heating phase and downflows during the cooling phase. This is consistent with the picture of chromospheric evaporation in the early phase and then coronal condensation in the decay phase of the flare, as reported in observations \\citep{anto83,anto85,anto90,schm87,ding96,bros03, raft09} and 1D hydrodynamic models \\citep{dosc82, fish85a, fish85b}. Therefore, the EBTEL computed flow can be also compared with Doppler shift measurements to provide another constraint to flare heating models.  \\begin{table}[h] \\centering \\tiny \\caption{\\textsc{\\small{Comparison among the studies referring to the EBTEL model and flare heating}}} \\begin{tabular}{cccc} \\tableline \\tableline \\multicolumn{1}{c}{studies} & heating function    & loop  number    & compare with observations \\\\ &                                  &                            & (instruments) \\\\ \\tableline \\tableline \\cite{raft09}    & inferred from hard X-ray light curve                & 1 & temperature, emission measure, velocity\\\\ &                                                                             &    & (\\textit{RHESSI}, \\textit{GOES}, \\textit{TRACE}, CDS)  \\\\ \\tableline \\cite{hock12} & free parameters                                                 & 22 & EUV light curves\\\\ &                                                                             &      & (EVE)\\\\ \\tableline \\cite{qiuj12}    & inferred from spatially resolved UV light curves    & thousands & X-ray and EUV light curves\\\\ &                                                                                   &                    & (\\textit{GOES}, AIA) \\\\ \\tableline this study       & inferred from spatially resolved UV light curves     & 2                  & EUV light curves, velocity\\\\ &                                                                                    &                     &(AIA, EIS)  \\\\ \\tableline \\end{tabular} \\label{comtab} \\end{table}  Following \\cite{qiuj12}, we will observationally infer heating functions from UV light curves and compute plasma evolution in individual flare loops. It, however, takes a few steps ahead of previous studies by making use of the most stringent observational constraints, in both the input heating functions and the output radiation and flow signatures. The novel aspects in this study include (1) the heating functions and predicted coronal radiations are derived and compared for different loops that are identified from imaging EUV observations and exhibit distinctive evolution patterns; (2) both radiation and velocity signatures are compared with observations in multiple lines by multiple instruments, AIA and the EUV Imaging Spectrometer (EIS; \\citealt{culh07}) on board \\textit{Hinode}. Such more comprehensive analysis and modeling allow us to better constrain the loop heating model and diagnose flare heating more accurately. Whereas \\cite{qiuj12} have used EBTEL to model a few thousand loops with the same set of free parameters, in this study, we will focus on modeling and analyzing only two loops in order to refine determination of heating functions. The 0D EBTEL allows us to investigate more efficiently variations of these free parameters and gain insight into the underlying physics that governs these parameters. Among these parameters, one is related to the heating rate, and the other describes the energy loss through the lower atmosphere. These parameters might not be the same in different loop systems, and will be examined in the present study. It also should be noted that the 0D EBTEL code is bench-marked by more sophisticated 1D hydrodynamical simulations with reasonable agreement \\citep{klim08}; on the other hand, with the EBTEL model combined with our method, we will be able to determine the first-order heating rates to be used as a valuable reference for more sophisticated 1D hydrodynamic models. In the following Section \\ref{obs}, we describe flare observations. The model computation and comparison with observations are presented in Section \\ref{result}.  We give the summary and discussions in Section \\ref{discon}. ",
        "conclusions": "\\label{discon}  We analyze and model two distinctive loop systems that are treated as two big loops in an M1.0 flare. Using the empirical heating functions inferred from observed UV light curves, we compute the evolution of plasmas in these loops via EBTEL. With the model result we reproduce the EUV synthetic fluxes and compare them with observations by AIA and EIS for each of the two loops. We also compare the computed mean enthalpy flow velocity with the measured Doppler velocity from EIS. The model results agree with observations in many aspects. The two different loops show different coronal plasma evolution patterns, with the short loop evolving more quickly than the long loop. These patterns are revealed in observations and as well reproduced in the model. Our results suggest that the observationally constrained heating rates and the EBTEL modeling with only two free parameters provide a reasonable description of the mean properties of heating and evolution of flare loops.  In our model, the loop length, and the timing and duration of the heating rate are derived directly from observations. Previous study using the similar method uses exactly the same set of parameters to scale the heating rate and energy loss rate for a few thousand loops. Our experiment of modeling two spatially resolved loop systems then allows us to investigate the variations of these model parameters  $\\lambda$ and $c_1$ in different loops. The parameter $\\lambda$ relates the UV emission counts directly from data to the amount of direct heating in the corona ($QL$), and $c_1$ scales the total energy loss through the transition region, or the base of the flare loop. Our results show that, to best match the mean properties of loops, different parameters have to be used. The best-match $\\lambda$ for the long loop is four times that of the short loop (see Table \\ref{partab}). On the other hand, $c_1$ for the long loop is one third that of the short loop. For the long loop, the value of $c_1$ is comparable with the average value determined in \\cite{qiuj12} who studied a few thousand long loops; and for the short loop, the optimal $c_1$ value is substantially larger. These differences may indicate different heating mechanisms and their subsequent low-atmosphere responses in the two loops. A plausible explanation for these differences may involve heating by electron beams. Imaging observations by {\\em RHESSI} suggest that during the early impulsive phase, there are hard X-ray sources located at the feet of the short loop. {\\em In-situ} heating by non-thermal electrons produces significant UV emission at these locations (\\citealt{allr05, cheng10, cheng12}, and references therein). Therefore, less amount of UV emission is contributed by direct heating in the corona, which heats the lower atmosphere through thermal conduction during the impulsive phase. A semi-quantitative analysis including non-thermal heating term in the EBTEL modeling is given by \\cite{liuw12}. A larger $c_1$ value used for the short loop also indicates a greater amount of radiative loss through the transition region than in long loops. This suggests that $c_1$ might not be an independent constant but may vary with the loop dynamics \\citep{carg12}.  Flow velocity is an important property in studying the dynamic evolution of flares. Spectroscopic velocity measurements can help diagnose the process of energy release and plasma evolution. During the impulsive phase, chromospheric evaporation (upflow), driven either by electron beams or thermal conduction \\citep{acto82, fish85a, fish85b, bros04, mill08}, provides indirect heating apart from the direct heating ($Q$) in the corona. In the cooling phase, coronal draining (downflow) dominates. In our present study, the spectroscopic observations do not adequately cover the impulsive phase of the flare, but our analysis shows that the modeled timing of the turn-over from upflow to downflow is consistent with the measured Doppler velocity in the decay phase, which supports the result of \\cite{brad10a,brad10b}. Our analysis shows the potential of using Doppler shift measurements to constrain the EBTEL model that invokes enthalpy flow in the flare evolution. In our future investigation, the high efficiency of the EBTEL model will allow us to model a good number of spatially resolved loops, and statistically investigate the relationship between energy release and plasma flow.  Despite its advantages, the EBTEL model does not reproduce the temperature distribution along the loop, which may explain the insufficient low temperature plasma emission from the model, as compared with observations. Furthermore, recent studies suggest that the observed evaporation velocity during the impulsive phase is temperature dependent \\citep{mill09, chen10, ying11}, which requires more sophisticated models, such as 1D hydrodynamic model, to compute the plasma evolution especially during the impulsive heating phase. The temporal properties of heating rates determined in this study can be used as the reference for 1D models; on the other hand, a 1D model includes more assumptions or free parameters such as those describing distribution of heating along the loop and therefore invokes addititional observational constraints."
    },
    "1208/1208.0265_arXiv.txt": {
        "abstract": "{The geometry of the accretion flow around stellar mass and supermassive black holes depends on the accretion rate. Broad iron emission lines originating from the irradiation of cool matter can indicate that there is an inner disk below a hot coronal flow.} {These emission lines have been detected in X-ray binaries. Observations with the {\\it{Chandra X-ray Observatory, XMM Newton}} {\\rm {and}} {\\it{Suzaku}} {\\rm{ }} have confirmed the presence of these emission lines also in a large fraction of Seyfert-1 active galactic nuclei (AGN). We investigate the accretion flow geometry for which broad iron emission lines can arise in hard and soft spectral state.} {We study an ADAF-type coronal flow,  where the ions are viscously heated and electrons receive their heat only by collisions from the ions and are Compton cooled by photons from an underlying cool disk.} {For a strong mass flow in the disk and the resulting strong Compton cooling only a very weak coronal flow is possible. This limitation allows the formation of ADAF-type coronae above weak inner disks in the hard state, but almost rules them out in the soft state.} {The observed hard X-ray luminosity in the soft state, of up to 10\\% or more of the total flux, indicates that there is a heating process that directly accelerates the electrons. This might point to the action of magnetic flares of disk magnetic fields reaching into the corona. Such flares have also been proposed by observations of the spectra of X-ray black hole binaries without a thermal cut-off around 200 keV.} ",
        "introduction": "With the advent of {\\it{XMM-Newton}} observations it has become obvious that broad iron emission lines are common in Seyfert galaxies (Nandra et al. 2007). The number of sources with detections of strong disk lines is growing with the increasing number of new observations by {\\it{XMM-Newton}}, {\\it{Chandra}} and {\\it{Suzaku}}. These Fe K-shell emission lines are the strongest X-ray emission lines in most AGN and X-ray binaries (see the review by Miller 2007). Averaging AGN X-ray spectra from deep {\\it{Chandra}} fields Falocco et al. (2011) significantly detected the broad component of the Fe line in the X-ray spectra of low luminosity AGN at low redshift.  Since iron emission lines can be understood as the reaction of an accretion disk to irradiation by an external source of hard X-rays, the obvious question is under which conditions the accretion geometry contains an inner disk and above it a hot coronal flow/ADAF. An inner disk is necessary to explain the relativistic lines. As commonly accepted, the mass accretion rate determines whether an optically thick and geometrically thin disk reaches inward to the innermost stable circular orbit (ISCO) or the disk is truncated and replaced by an advection-dominated optically thin hot flow in the inner region. A detailed definition of these states is given by Remillard and McClintock (2006). For very high accretion rates, a third, so-called steep-power law state can exist, which is observed during some outbursts of X-ray binaries. Though in principle the accretion flow geometry is the same for X-ray binaries and AGN (Narayan et al. 1998), it is easier to distinguish between different spectral states for galactic sources than for AGN. Walton et al. (2012) analyzed high quality spectra of the stellar-mass black hole binary XTE J1650-500 and the active galaxy MCG-6-30-15 and argued that the similarity of broad iron line features shows that the excess in both sources is of the same nature.  During the hard/low state, the power-law flux from the ADAF in the inner region is the main source of radiation. If the disk is truncated, no relativistic iron emission lines should be produced. However, from the earliest studies a broad iron emission line has been detected for Cygnus X-1 in low state ( Barr et al. 1985). Now {\\it{Chandra}} and {\\it{XMM-Newton}} observations have firmly established the existence of a relativistic Fe K emission line in spectra of several X-ray binaries. Reis et al. (2010) analyzed the X-ray spectra of black hole binaries in the canonical low/hard state and found that the accretion disk reaches inward to the ISCO. These detections were made mostly during outburst decline and can theoretically be understood in terms of re-condensation of gas into an inner disk (Meyer et al. 2007) in the hard spectral state, as a transient phenomenon following the state transition. This accretion flow geometry, an outer truncated disk, and inside only a weak inner disk and above it the main flow via a corona/ADAF, is the situation where broad iron emission lines are expected. There may be a gap between the outer standard disk and such a weak inner disk.  During the high/soft (thermal) state, the flux is dominated by the radiation from the inner disk. The observations of X-ray binaries have identified an additional hard power-law flux, sometimes with a non-thermal tail, the power-law flux that is evidence of the co-existence of disk and hot flow. In this case then, one has to keep in mind that heat conduction and the possible exchange of mass and angular momentum between disk and hot flow can be significant. Investigations of the spectra of X-ray binaries, especially Cyg X-1 (Ibragimov et al. 2005), led to the conclusion that magnetic flares are needed to explain the non-thermal radiation.Hints to magnetic fields also come from accretion disk winds, which are ubiquitous in black hole binaries in the soft state, as pointed out by Ponti et al. (2012). Miller et al. (2008) modeled spectra observed for GRO J1655-4 and concluded that these disk winds  are probably related to magnetic fields.  The interpretation of spectra is less clear for AGN. Despite the large number of observations, it is difficult to classify the source as in hard or soft state. Vasudevan \\& Fabian (2009) presented spectral  energy distributions for a sample of AGN, based on contemporaneous optical, ultraviolet, and X-ray observations, and determined the ratio of bolometric to Eddington luminosity, using the black hole mass measurements of Peterson et al. (2004). Tanaka et al. (1995) first observed an asymmetric disk line profile in the Seyfert-1 AGN MCG-6-30-15, which remains a very important source for studies of relativistic disk lines. If a relativistic line is detected, one might conclude that an inner disk exists, which is either (1) in the soft state, a standard disk inward to the ISCO with irradiation from a coronal layer or (2) in the hard state, only a very weak inner disk below the standard ADAF. The latter case seems to occur in X-ray binaries as a transient phenomenon of the intermediate spectral state. One might expect to find  the same accretion flow geometry for AGN. Meyer-Hofmeister \\& Meyer (2011) investigated the situation using Seyfert galaxies for which emission lines were clearly present (Nandra et al. 2007, Miller 2007), but found that some of these sources were apparently in the soft rather than the hard state.  That broad iron emission lines are common in Seyfert AGN suggests that they are not a transient phenomenon and that the accretion geometry then corresponds to a soft state disk with a hot coronal layer/ADAF above, at least for some of the observed sources. In this paper we investigate the structure of the hot coronal flow above a standard disk under the influence of Compton cooling. In Sect.2, we give a short summary of broad iron emission lines detected in X-ray binaries and AGN. Sect.3 concerns the physics of interaction of disk and corona around supermassive and stellar-size black holes. We determine the limitation of the coronal flow caused by Compton cooling (Sect.4). The restriction in the case of a strong disk mass flow is so low that the observed hard power-law flux in a number of sources requires an additional heating process, e.g. one produced by magnetic flares. This is consistent with the conclusions derived by Ibragimov et al. (2005) from fitting  X-ray binary spectra with a non-thermal tail and no cut-off at energies between 100 keV and 200 keV (Sect.5). Whether the situation is similar for AGN is difficult to judge since for these sources, a non-thermal power-law flux does not seem to be indicated, but cannot be not generally excluded. In Sect.6, we refer to observations for X-ray binaries that detect evidence of non-thermal Comptonization. A discussion of questions related to our results follows in Sect. 6. ",
        "conclusions": ""
    },
    "1208/1208.1780_arXiv.txt": {
        "abstract": "We present the results of our follow-up observation program of $\\gamma$-ray sources detected by the Large Area Telescope (LAT) on board the {\\it Fermi Gamma-ray Space Telescope}. 26 blazars and 39 sources unidentified at other wavelengths were targeted at IRSF 1.4 m telescope equipped with the SIRIUS/SIRPOL imager and polarimeter. $H$-band magnitudes of the blazars at the epoch of 2010 Dec -- 2011 Feb are presented, which reveal clear flux variation since the Two Micron All Sky Survey observations and can be useful data for variation analyses of these objects in longer periods. We also find that nearly half of the $\\gamma$-ray blazars are highly ($>$10 \\%) polarized in near-infrared wavelengths. Combining the polarization and variation properties, most ($\\sim$90 \\%) of the blazars are clearly distinguished from all other types of objects at high Galactic latitudes. On the other hand, we find only one highly polarized and/or variable object in the fields of unidentified sources. This object is a counterpart of the optical variable source PQV1 J131553.00$-$073302.0 and the radio source NVSS J131552$-$073301, and is a promising candidate of new $\\gamma$-ray blazars. From the measured polarization and variation statistics, we conclude that most of the {\\it Fermi}/LAT unidentified sources are not likely similar types of objects to the known $\\gamma$-ray blazars. ",
        "introduction": "A new era of blazar studies has arrived with the advent of the {\\it Fermi Gamma-ray Space Telescope} ({\\it Fermi}). {\\it Fermi} has been carrying out an all-sky survey with its main instrument Large Area Telescope (LAT) since the science mission phase started in 2008. The first {\\it Fermi}-LAT catalog \\citep[1FGL;][]{{2010ApJS..188..405A}} lists 1451 $\\gamma$-ray detections, in which 821 sources are associated (or identified) with objects found at other wavelengths. The majority of the associated objects are active galactic nuclei (AGNs) dominated by blazars, while other extragalactic sources as well as Galactic objects such as pulsars and supernova remnants make smaller contributions. The remaining 630 sources are left unidentified in 1FGL.   Blazars are considered to be AGNs whose jets are aligned with the observer's line-of-sight. Emission from the relativistically boosted jets dominates observed flux, resulting in two broad peaks in the spectral energy distribution (SED); one at radio to X-rays arising from synchrotron emission of accelerated, high-energy particles, and another at X-ray to $\\gamma$-ray arising from the inverse Compton scattering of the lower energy photons. Because of these emission mechanisms, blazars are characterized by strong radio, X-ray, and $\\gamma$-ray radiations as well as high polarization and variation across the entire SED. Since blazars dominate the $\\gamma$-ray sky at high Galactic latitudes, the {\\it Fermi} all-sky survey is expected to shed new light on this relatively rare and poorly understood population.   An investigation of blazar radiation mechanisms (hence the intrinsic SED) is important not only for revealing the nature of blazars themselves, but also for measuring the extragalactic background light (EBL). Since high-energy photons from blazars interact with optical to near-infrared (IR) EBL in the intergalactic space, observations of distant blazars can be used to infer the EBL spectrum if the intrinsic blazar SED is precisely known. With this indirect method, \\citet{aharonian06} obtained the significantly-lower upper limits of near-IR EBL than those derived from the direct measurements by, e.g., \\citet{matsumoto05}. Recently \\citet{matsuoka11a} succeeded in a direct measurement of optical EBL by re-analyzing the {\\it Pioneer 10/11} data, which is on the smooth extension of the near-IR upper limits obtained by \\citet{aharonian06}.   In this paper, we present the results of our near-IR follow-up observation program of $\\gamma$-ray blazars and unidentified sources in 1FGL. Despite of the wealth of information there, near-IR wavelength of the whole AGN population is still poorly understood \\citep[e.g.,][]{matsuoka07,matsuoka08,matsuoka11b}. We aim to quantify the most distinct features of blazars in near-IR wavelengths, i.e., polarization and variation, in part as a benchmark for future observations. At the same time, we explore the nature of 1FGL unidentified sources by comparing their near-IR properties to the known $\\gamma$-ray blazars. Dominance of blazars in the $\\gamma$-ray sky implies that some of the unidentified sources at high Galactic latitudes are similar objects missed in the past surveys at other wavelengths. Revealing the origin of these unidentified $\\gamma$-ray emissions is of greatest importance, therefore many follow-up studies are being dedicated to this subject \\citep[e.g.,][]{ackermann11b}. We aim to provide complementary information to these follow-up programs from a near-IR view point.   While the second {\\it Fermi}-LAT catalog \\citep[2FGL;][]{nolan12} has already been released, this paper is based on 1FGL for consistency with the sample selection of presented observations. We discuss the new 2FGL identifications of the sample later. Magnitudes are presented on the Vega-based system throughout this paper. ",
        "conclusions": "We present the results of photometry and polarimetry measurements for the blazars in Table \\ref{tab:obs_bzx}. Their polarization and variation since 2MASS observations, as well as those for all other sources detected in the blazar fields, are plotted in Figure \\ref{var_stat_bzx}. Strong variability of the blazars is evident; 22 out of 26 blazars (85 \\%) have $| H_{\\rm IRSF} - H_{\\rm 2MASS} | > 0.25$ mag while only one of the other detected sources shows such variation ($\\sim$0.3 mag; it is likely a contaminating normal star). The blazars are also characterized by pronounced polarization; 10 out of 25 blazars (40 \\%; polarization was not measured for a blazar associated to J0038.4$-$2504 due to a large ($>$ 0.05 mag) photometry error) have $P >$ 10 \\%, while all other types of objects are much less polarized. The presented fraction, 40 \\%, represents a useful benchmark for planning future optical/near-IR polarization follow-up programs of $\\gamma$-ray selected blazars.  \\begin{figure} \\epsscale{1.0} \\plotone{var_stat_bzx.eps} \\caption{Polarization and variation of 326 sources detected in the blazar fields. The diamonds represent the blazars associated to 1FGL sources while the dots represent all other objects. The arrows denote upper limits of polarization. \\label{var_stat_bzx}} \\end{figure}   From the above results, we can derive an expected number of blazars which should be discovered by their high polarization and/or variation in the fields of unidentified sources {\\it if they are similar but unknown blazars}. In total, 23 out of 26 blazars (88 \\%) have pronounced polarization ($P > 10$ \\%) or variation ($| H_{\\rm IRSF} - H_{\\rm 2MASS} | > 0.25$ mag). Since 39 unidentified sources were observed, we would have $\\sim$35 objects with such distinguishable properties under the above assumption. However, $\\sim$35 \\% of them would be fainter than our limiting magnitude $H_{\\rm IRSF}^{\\rm lim} \\simeq 16.0$ mag based on the 2MASS magnitude distribution of 1FGL blazars. Furthermore, $\\sim$32 \\% of them would be outside SIRPOL field-of-view (7'.7 $\\times$ 7'.7 arcmin$^2$) when $\\gamma$-ray positions are used as the telescope pointing center, based on the distribution of distance between $\\gamma$-ray positions and associated blazars. Considering these restrictions, we expect to find 15 blazars out of 39 fields of unidentified sources if their $\\gamma$-ray is indeed emitted by a similar population to 1FGL blazars.   Figure \\ref{var_stat_una} shows polarization and variation of 608 objects detected in all the fields of unidentified sources. While three objects are found to have high variation ($| H_{\\rm IRSF} - H_{\\rm 2MASS} | > 0.25$ mag), two of them with $P < 7$ \\% and $| H_{\\rm IRSF} - H_{\\rm 2MASS} | \\sim 0.3$ mag are likely contaminations; it is consistent with one contamination out of 326 objects in the blazar fields (Figure \\ref{var_stat_bzx}). On the other hand, the object at $P \\sim 9$ \\% and $| H_{\\rm IRSF} - H_{\\rm 2MASS} | \\sim 1.2$ mag is a promising candidate of new $\\gamma$-ray blazars. It is found in J1315.6$-$0729 field, at R.A. 13$^{\\rm h}$15$^{\\rm m}$52$^{\\rm s}$.98, Dec. $-$07$^{\\rm d}$33$^{\\rm m}$01$^{\\rm s}$.99 (J2000.0) with $H$-band brightness $H_{\\rm IRSF} = 14.1$ mag and $H_{\\rm 2MASS} = 15.3$ mag. Its counterparts are found in the NED\\footnote{ The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.}; the optical variable source PQV1 J131553.00$-$073302.0 \\citep{bauer09} and the radio source NVSS J131552$-$073301. We plan to carry out a spectroscopic follow-up observation of this object in the near future.  \\begin{figure} \\epsscale{1.0} \\plotone{var_stat_una.eps} \\caption{Same as Figure \\ref{var_stat_bzx}, but for the fields of unidentified sources. \\label{var_stat_una}} \\end{figure}   Except for the above blazar candidate, we found no clear sign of variable or polarized objects in the fields of unidentified sources. The apparent inconsistency between the expected number of blazar candidates as estimated above (15) and the actual number (1) indicates that $\\gamma$-ray emission of the unidentified sources arises from other types of objects than known 1FGL blazars. They could be non-blazar active galaxies, starburst galaxies, or Galactic objects such as pulsars and supernova remnants at relatively high Galactic latitudes, as well as active galaxies with jets but without strong polarization and variability at near-IR wavelengths. In this regard, it is noteworthy that some of them are already identified (or associated) in the latest 2FGL catalog and the {\\it Fermi}-LAT second AGN catalog \\citep[2LAC;][]{ackermann11} as summarized in Table \\ref{obslist_id_by_2fgl}. Many of them are pulsars, which is consistent with our observation results. While J2330.3$-$4745 is associated to the blazar PKS 2326$-$477, their separation is relatively large and the blazar is outside the field-of-view of our SIRPOL observation.        \\begin{table*} \\caption{2FGL identification (association) of 1FGL \"unidentified\" sources in our sample}\\label{obslist_id_by_2fgl} \\begin{center} \\begin{tabular}{lll} \\hline 1FGL ID        & Associated source      & Object type  \\\\ \\hline J0001.9$-$4158 & 1RXS J000135.5$-$41551 & active galaxy of uncertain type\\\\ J0101.0$-$6423 & PSR J0101$-$6422       & pulsar \\\\ J0223.0$-$1118 & 1RXS J022314.6$-$11174 & active galaxy of uncertain type\\\\ J0335.5$-$4501 & IRXS J033514.5$-$44592 & active galaxy of uncertain type\\\\ J0614.1$-$3328 & PSR J0614$-$3330       & pulsar \\\\ J1124.4$-$3654 & PSR J1124$-$36         & pulsar \\\\ J1141.8$-$1403 & 1RXS J114142.2$-$14075 & active galaxy of uncertain type\\\\ J1231.1$-$1410 & PSR J1231$-$1411       & pulsar \\\\ J1312.6$+$0048 & PSR J1312$+$00         & pulsar \\\\ J2241.9$-$5236 & PSR J2241$-$5236       & pulsar \\\\ J2330.3$-$4745 & PKS 2326$-$477         & blazar \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}"
    },
    "1208/1208.1263_arXiv.txt": {
        "abstract": "We combine $VI$ photometry from OGLE-III with $VVV$ and 2MASS measurements of $E(J-K_{s})$ to resolve the longstanding problem of the non-standard optical extinction toward the Galactic bulge. We show that the extinction is well-fit by the relation $A_{I} = 0.7465{\\times}E(V-I) + 1.3700{\\times}E(J-K_{s})$, or, equivalently, $A_{I} = 1.217{\\times}E(V-I)(1+1.126{\\times}(E(J-K_{s})/E(V-I)-0.3433))$.  The optical and near-IR reddening law toward the inner Galaxy approximately follows an $R_{V} \\approx 2.5$ extinction curve with a dispersion ${\\sigma}_{R_{V}} \\approx 0.2$, consistent with extragalactic investigations of the hosts of type Ia SNe. Differential reddening is shown to be significant on scales as small as as our mean field size of 6$\\arcmin$. The intrinsic luminosity parameters of the Galactic bulge red clump (RC) are derived to be $(M_{I,RC}, \\sigma_{I,RC,0}, (V-I)_{RC,0}, \\sigma_{(V-I)_{RC}}, (J-K_{s})_{RC,0}) = (-0.12, 0.09, 1.06, 0.121, \\newline 0.66)$. Our measurements of the RC brightness, brightness dispersion and number counts allow us to estimate several Galactic bulge structural parameters. We estimate a distance to the Galactic center of 8.20 kpc. We measure an upper bound on the tilt $\\alpha \\approx 40^{\\circ}$ between the bulge's major axis and the Sun-Galactic center line of sight, though our brightness peaks are consistent with predictions of an N-body model oriented at $\\alpha \\approx 25^{\\circ}$. The number of RC stars suggests a total stellar mass for the Galactic bulge of $\\sim2.3{\\times}10^{10} M_{\\odot}$ if one assumes a canonical Salpeter IMF, or  $\\sim1.6{\\times}10^{10} M_{\\odot}$ if one assumes a bottom-light Zoccali IMF. ",
        "introduction": "\\label{sec:Introduction} The central bulge of the Milky Way Galaxy is the only stellar spheroid for which we can measure detailed abundances, ages and all six phase space dimensions for individual stars, as well as the luminosity function and spatial distribution for the population as a whole. Some $\\sim$10\\% of the Milky Way's stars are bulge stars, including a disproportionate number of the oldest and most metal-rich stars. It is therefore evident that any theory of Galaxy formation and evolution is required to reproduce the observed properties of the bulge, and conversely, that the properties of the bulge should be measured as precisely and accurately as possible to best discriminate between different Galaxy formation models.  However, as scientifically desirable as this greater project may be, it is also difficult due to several challenges that prevent further, deeper understanding of the bulge\\footnote{Henceforth, we almost exclusively refer to the bar/bulge of the Milky Way as the bulge for the sake of consistent representation. We do recognise that these two words have different meanings, but the kinematic decomposition of the Galaxy's central population remains a matter of active investigation and controversy at this time.}. There are significant correlations between chemistry and kinematics \\citep{2010A&A...519A..77B,2012A&A...546A..57U,2012ApJ...756...22N}, distinct chemical subgroups \\citep{2012arXiv1205.4715A}, gradients in metallicity \\citep{2008A&A...486..177Z}, deviations from the classical picture of the triaxial ellipsoid at large separations from the minor axis \\citep{2005ApJ...630L.149B,2008A&A...491..781C}, and both large  \\citep{2010ApJ...721L..28N,2010ApJ...724.1491M} and small \\citep{2001A&A...379L..44A,2005ApJ...621L.105N} separations from the plane. These issues necessitate larger data sets and better models. The viewing angle $\\alpha$ between the bulge's major axis and the sun-Galactic center line of sight remains undetermined, with best-fit values ranging from from $\\alpha=13^{\\circ}$ \\citep{2007A&A...465..825C}  to $\\alpha=44^{\\circ}$ \\citep{2005ApJ...630L.149B}. This prevents further understanding of the inner Galaxy's gravitational potential, as uncertainties in the value of $\\alpha$ are degenerate with those of the bulge's axis ratios \\citep{1997ApJ...477..163S} and rotation speed \\citep{2010ApJ...720L..72S}.  The most significant source of uncertainty, however, is the extinction. It averages $A_{K} \\approx 3$ toward the  Galactic center \\citep{2010A&A...511A..18S}, suggesting $A_{V} \\approx 50$ \\citep{2008ApJ...680.1174N}. For most of the bulge, values of $A_{V} = 2$ are typical \\citep{2004MNRAS.349..193S}. The high values of reddening close to the plane render it quite difficult to obtain spectroscopic observations, proper motions, and stellar density maps. Further from the plane, these can be obtained, but not fully understood due to significant zero-point uncertainties in the extinction, and indirectly, the distance.  A further complication arises from the fact the extinction toward the inner Galaxy is not only large but also non-standard. This was first suggested by \\citet{2000ApJ...528L...9P} as a solution to the anomalous colors of bulge RR Lyrae \\citep{1999ApJ...521..206S} and red clump (RC) stars \\citep{1999AcA....49..319P}.  \\citet{2001ApJ...547..590G} and \\citet{2003ApJ...590..284U} were the first to demonstrate that the reddening law toward the inner Galaxy is described by smaller total-to-selective ratios than the ``standard'' values measured for the local interstellar medium, implying a steeper extinction curve and thus a smaller characteristic size for dust grains \\citep{2003ARA&A..41..241D}. \\citet{2003ApJ...590..284U} measured $ dA_{I}/dE(V-I) \\approx 1.1$ (denoted  ${\\Delta}A_{I}/{\\Delta}E(V-I)$ in that work) toward several bulge fields, much smaller than the value of $dA_{I}/dE(V-I) \\approx 1.45$ suggested by the standard interstellar extinction curve of $R_{V}=3.1$ \\citep{1989ApJ...345..245C,1994ApJ...422..158O}.  \\citet{2003ApJ...590..284U} showed that applying the same methodology to observations of the Large Magellanic Cloud taken with the same instruments yielded $dA_{I}/dE(V-I) \\approx 1.44$, the standard value, demonstrating the robustness of the result. Further, not only was the reddening law toward the bulge found to be non-standard, it was also found to be rapidly varying between sightlines, with values of $dA_{I}/dE(V-I)$ ranging from $0.94\\pm0.01$ to $1.16\\pm0.03$. The steeper extinction law toward the inner Galaxy has been subsequently confirmed with observations using \\textit{Hubble Space Telescope (HST)} optical filters by \\citet{2010A&A...515A..49R}, by analysis of RR Lyrae stars in OGLE-III \\citep{2012ApJ...750..169P}, and also in the near-IR \\citep{2008ApJ...680.1174N,2009MNRAS.394.2247G,2010A&A...511A..18S}. Meanwhile, \\citet{2009ApJ...707..510Z} and \\citet{2011ApJ...737...73F} both found that the extinction law toward the inner Galaxy was \\textit{shallower} (greyer) in the mid-IR.  The variations in both the extinction and the extinction law made it difficult to reliably trace the spatial structure of the bulge \\citep{2010AcA....60...55M}. Applying the $VI$ extinction maps of \\citet{2004MNRAS.349..193S} to the bulge color-magnitude diagram (CMD) implied a distance to the Galactic center of $\\sim$9 kpc \\citep{2007MNRAS.378.1064R,2009A&A...498...95V}, a large value relative to the geometrically determined distances to the Galactic center of $7.62 \\pm 0.32$ kpc \\citep{2005ApJ...628..246E}, $8.27 \\pm 0.29$ kpc \\citep{2012arXiv1207.3079S},  and $8.4 \\pm 0.4$ kpc \\citep{2008ApJ...689.1044G}. As the structure of the inner Galaxy is a very sensitive probe of the environmental conditions in which the Galaxy formed and evolved \\citep{2005MNRAS.358.1477A,2012MNRAS.422.1902I}, an accurate spatial determination of the bulge's morphology would yield a powerful test of Galaxy formation models. Moreover, investigations of the metallicity distribution function of bulge giants have had to rely on imprecise and potentially inaccurate estimates of surface gravity and photometric temperature, further reducing our ability to probe the primordial conditions of the Galaxy.  We resolve these issues in this investigation by combining OGLE-III observations in $V$ and $I$ with $VVV$ and 2MASS measurements of $E(J-K_{s})$ \\citep{2012A&A...543A..13G}. We confirm previous findings that the $VI$ extinction toward the inner Galaxy is steeper than standard, but also show that it is a little less steep than previously assumed. We show that this is likely due to an effect we label ``composite extinction bias'', which makes it unphysical to extrapolate a slope of $dA_{I}/dE(V-I)$ to $E(V-I)=0$. Our parameterization for the extinction,  $A_{I} = 0.7465{\\times}E(V-I) + 1.3700{\\times}E(J-K_{s})$, is less sensitive to the rapid variations in the extinction law than the computation of slopes $dA_{I}/dE(V-I)$. Whereas the latter must be computed from an ensemble of measurements spread across 30$\\arcmin$ or more, the former can be directly measured for each $\\sim 6\\arcmin \\times 6\\arcmin$ sightline.  The structure of this paper is as follows. We summarize the data used in Section \\ref{sec:Data}. Our methodology for measuring the parameters of the RC is described in Section \\ref{sec:RCmeasurements}, and we derive  the intrinsic RC luminosity parameters in Section \\ref{sec:Calibration}. We briefly state the properties of reddening that would be expected using a standard reddening curve in Section \\ref{Sec:ReddeningTheory}. The reddening measurements are presented and discussed in Section \\ref{sec:Reddening}, including comparisons to the reddening maps of Schlegel, Finkbeiner \\& Davis (SFD, \\citealt{1998ApJ...500..525S}) and the derivation of an empirical rule to estimate differential reddening. Methods to convert the reddening into an extinction using a single color are demonstrated to inevitably fail in Section \\ref{sec:ReddeningLawEstimates}, and a more successful extinction law is derived in Section \\ref{sec:RJKVI} by including information from both $E(V-I)$ and $E(J-K_{s})$. We demonstrate that reddening constraints from MACHO photometry may have been misinterpreted in Section \\ref{sec:MACHO}. In section \\ref{sec:GalacticStructure}, we show that our dereddened apparent magnitudes suggest a distance to the Galactic center $R_{0}=8.20$ kpc, and a tilt between the Galactic bulge's major axis and the sun-Galactic center line of sight no greater than $\\alpha \\approx 41^{\\circ}$. We translate these measurements into constraints for microlensing events toward the bulge in Section \\ref{sec:MicrolensingConstraint}. We analyze our number counts for the RC in Section \\ref{sec:GalacticStructure2} and show that combining these with the assumptions of standard stellar evolution and a Salpeter IMF yields an estimated Galactic bulge stellar mass of $M \\sim 2.3{\\times}10^{10} M_{\\odot}$. The thickness of the Galactic bulge is discussed in Section \\ref{sec:GalacticStructure3}. The data structure is summarized in Section \\ref{Sec:DataSummary}. Results are discussed in Section \\ref{sec:Discussion}. ",
        "conclusions": "\\label{sec:Discussion} Our solution to the observational problem of the non-standard $VI$ extinction toward the inner Galaxy mitigates what has been one of the dominant sources of uncertainty in studies of the Galactic bulge. The extinction law is on average steeper in the optical with $\\sim$30\\% variations superimposed, and is well-fit by the relation $A_{I} = 0.7465{\\times}E(V-I) + 1.3700{\\times}E(J-K_{s})$. The residuals to the extinction fit is now reduced to no more than 0.06 mag, and the estimate of $R_{0} = 8.20$ kpc is consistent with there being no bias in our fit to the extinction law.  In the course of making these measurements, we have also measured that differential reddening averages ($\\sim$9\\% of total reddening for small fields), and the intrinsic luminosity parameters for the bulge RC. These will be of use to future bulge studies.  The mean value of $A_{I}/E(V-I)=1.217$ suggests $R_{V}=2.5$, the mean value of $R_{JKVI} = 0.3433$ suggests $R_{V}=2.6$, and thus both the measurements investigated here are consistent with an $R_{V} \\approx 2.5$ extinction curve. Measuring the extinction curve in other bandpasses could potentially have major implications for cosmology. Our inferred extinction curve is consistent with the values of $R_{V} \\approx 2.5$  inferred in studies of the extinction toward extragalactic type Ia SNe by \\citet{2010A&A...523A...7G} and \\citet{2011A&A...529L...4C}. In particular, the hierarchical Bayesian analysis of \\citet{2011ApJ...731..120M} found that the extinction law toward SNe Ia went as $R_{V}=2.5-2.9$ for $A_{V} \\leq 0.4$, and steepened at higher extinctions. \\citet{1999ApJ...523..617F} also reported a range in the extinction laws of 23 lensed galaxies of $1.5 \\leq R_{V} \\leq 7.2$. Since none of the combined 51 measurements of \\citet{1989ApJ...345..245C} and \\citet{1994ApJ...422..158O} reach such low values of $R_{V}$, there is no reason to expect that extrapolation of this empirical law will behave adequately in domains that lie well beyond its calibration. Indeed, though the values of $A_{I}/E(V-I)=1.217$ and $R_{JKVI} = 0.3433$ are consistent with each other, the change in  $A_{I}/E(V-I)$ as $R_{JKVI}$ changes does not go at the rate predicted by \\citet{1989ApJ...345..245C} and \\citet{1994ApJ...422..158O}, as shown in Figure \\ref{Fig:LongMagMosaic3B}. This domain of the extinction law therefore warrants further investigation. Some constraints could be extracted by combining our results with the recent study of \\citet{2012ApJS..201...35N}, who measured reddening values for the color $([3.6{\\mu}] - [4.5{\\mu}])$.  Our measurements of the number counts, brightness dispersion, mean brightness and full error matrix thereof for $\\sim$9,000 RC centroids toward the bulge may be one of the most potent means for constraining the structural parameters of the Galactic bulge. In Sections \\ref{sec:GalacticStructure} and \\ref{sec:GalacticStructure2}, we have sketched how these data could be used to constrain the morphology and mass of the bulge, without going to the full formalism employed by \\citet{1995ApJ...445..716D}, \\citet{1997ApJ...477..163S} and \\citet{2007MNRAS.378.1064R}. Moreover, as the data have improved, it is now time for the models to improve as well. The use of N-body models by \\citet{2007MNRAS.378.1165R}, \\citet{2010ApJ...720L..72S}, \\citet{2011ApJ...734L..20M}, \\citet{2012ApJ...756...22N} and \\citet{2012ApJS..201...35N} are encouraging steps in that direction."
    },
    "1208/1208.6443_arXiv.txt": {
        "abstract": "We study magneto-elastic oscillations of highly magnetized neutron stars (magnetars) which have been proposed as an explanation for the quasi-periodic oscillations (QPOs) appearing in the decaying tail of the giant flares of soft gamma-ray repeaters (SGRs). We extend previous studies by investigating various magnetic field configurations, computing the Alfv\\'en spectrum in each case and performing magneto-elastic simulations for a selected number of models. By identifying the observed frequencies of $28\\,$Hz (SGR 1900+14) and $30\\,$Hz (SGR 1806-20) with the fundamental Alfv\\'en QPOs, we estimate the required surface magnetic field strength. For the magnetic field configurations investigated (dipole-like poloidal, mixed toroidal-poloidal with a dipole-like poloidal component and a toroidal field confined to the region of field lines closing inside the star, and for poloidal fields with an additional quadrupole-like component) the estimated dipole spin-down magnetic fields are between $8\\times 10^{14}$G and $4\\times10^{15}$G, in broad agreement with spin-down estimates for the SGR sources producing giant flares. A number of these models exhibit a rich Alfv\\'en continuum revealing new turning points which can produce QPOs. This allows one to explain most of the observed QPO frequencies as associated with magneto-elastic QPOs. In particular, we construct a possible configuration with two turning points in the spectrum which can explain all observed QPOs of SGR 1900+14. Finally, we find that magnetic field  configurations which are entirely confined in the crust (if the core is assumed to be a type I superconductor) are not favoured, due to difficulties in explaining the lowest observed QPO frequencies ($f\\lesssim30\\,$Hz). ",
        "introduction": "The discovery of quasi-periodic oscillations (QPOs) in the decaying tail of a giant flare of a soft-gamma ray repeater\\footnote{SGRs are assumed to be highly magnetized neutron stars showing repeated bursts in the soft-gamma ray spectrum \\citep{Duncan1992}} (SGR) by \\cite{Israel2005} \\citep[see also][and references therein]{Watts2007} has stimulated strong interest in the theoretical modelling of oscillations of magnetized neutron stars (magnetars). The observed frequencies of the QPOs are roughly $18$, $26$, $30$, $92$, $150$, $625$, and $1840\\,$Hz for the outburst of SGR 1806-20 and $28$, $53$, $84$, and $155\\,$Hz for SGR 1900+14, respectively. Since the first attempts to explain these QPOs in terms of torsional, purely shear modes of the solid crust of the neutron star \\citep{Duncan1998,Strohmayer2005,Piro2005, Sotani2007, Samuelsson2007, Steiner2009} the theoretical understanding of the oscillations has evolved. The most recent studies \\citep{Gabler2011letter,Gabler2012,Colaiuda2011, vanHoven2011,vanHoven2012} favour global magneto-elastic oscillations as an explanation for the lower frequency ($f<200\\,$Hz) QPOs. In these works the coupling through the magnetic field leads to a very efficient absorption of the shear modes of the crust into the core even for magnetic field strengths well below those estimated for magnetars, $B_{15}\\sim 1$ with $B_{15}\\equiv B[10^{15} \\,\\mathrm{G}] $ \\citep{Duncan1992}. Nonetheless, long-lived QPOs which have predominantly Alfv\\'en character have also been studied \\citep{Sotani2008,Cerda2009,Colaiuda2009}. The frequencies of successive overtones of these QPOs are found in integer relations and, thus, can potentially explain some of the observed frequency relations. In two previous papers~\\citep{Gabler2011letter,Gabler2012} we have shown that these oscillations can reach the surface of the star with significant amplitudes only for magnetic fields stronger than $B_{15}\\gtrsim 1$, if the internal field has a global, dipole-like structure.  A successful interpretation of the observed QPOs has the potential to constrain the equation of state (EoS) of the cold matter at supranuclear densities which occur only in the core of a neutron star \\citep{Samuelsson2007}. Therefore, neutron stars are a unique laboratory that can be used to increase our knowledge of the fundamental physics describing the interaction of nucleons and other elementary particles at the conditions inside those stars.  While the current models of magneto-elastic oscillations are promising to explain some of the observed QPO frequencies, all existing studies are restricted to a very limited set of  magnetic field configurations. The aim of the present work is to explore the influence of the magnetic field configuration in the magneto-elastic model of QPOs to assess its validity. It must be stressed that very little is known from observations about the internal magnetic field configurations of magnetars. The only observational constraint is the presence of a strong external field which is responsible for the observed spin-down of the magnetar. Both purely toroidal and purely poloidal magnetic fields are known to be unstable in unstratified stars, i.e. stars with a barotropic EoS which depends on one parameter only \\citep{Tayler1973,Markey1973}. Non-linear simulations confirm these instabilities \\citep[][]{Braithwaite2006, Kiuchi2011, Ciolfi2011, Lasky2011, Lander2011, Lander2011b} and suggest that a twisted torus configuration, with mixed poloidal and toroidal field, is expected. There have been some attempts to model magnetic field configurations as equilibrium axisymmetric configurations with mixed poloidal and toroidal field \\citep{Colaiuda2008,Kiuchi2008,Ciolfi2009, Lander2009}. However, no stable configuration has yet been found for barotropic stars \\citep{Lander2012b}. Nevertheless, different possibilities could stabilize magnetic fields in stars: a more complicated magnetic field structure (possibly non-axisymmetric), stratification, or the presence of a solid crust. Until this issue is settled, it is legitimate to explore the most general setup of magnetic field configurations possible and analyze their influence on the QPO frequencies, leaving aside  the stability of the configurations. This paper, in which we explore  a wide range of plausible axisymmetric magnetic field configurations, summarizes our findings in this direction.  We begin this paper in Section\\,\\ref{sec_theory} with a short overview of our theoretical and numerical framework, described in more detail in \\cite{Cerda2009} and \\cite{Gabler2012}. We also describe in detail in this section how to obtain different magnetic field configurations with a new implementation of the {\\small MAGSTAR} tool of the {\\small LORENE} library which we call {\\small MAGNETSTAR}. In the following sections we consider different magnetic field configurations, including dipole-like fields in Section\\,\\ref{sec_dipoles} and quadrupole-like fields in Section\\,\\ref{sec_quadru}, where  mixed quadrupole-dipole-like fields are also considered. We further divide the dipole-like fields into purely poloidal fields (Section\\,\\ref{sec_poloidal}) and mixed toroidal-poloidal fields (Section\\,\\ref{sec_field_LoreneII}). All magnetic field configurations considered up to Section\\,\\ref{sec_quadru} penetrate the whole volume of the star. In Section\\,\\ref{sec_crustfield} we study the case of a magnetic field confined to the crustal region, a situation motivated by the indications that neutron stars may contain  superconducting protons \\citep{Page2011,Shternin2011} in which case the magnetic field could be expelled from the core of the neutron star, if it is a type I superconductor. A summary of our results is presented in Section\\,\\ref{sec_conclusions}.  We use units where $c=G=1$ with $c$ and $G$ being the speed of light and the gravitational constant, respectively. Latin (Greek) indices run from 1 to 3 (0 to 3). Partial and covariant derivatives are indicated by a comma and a semicolon, respectively. We apply the Einstein summation convention. ",
        "conclusions": "\\label{sec_conclusions}  We studied the effect of different magnetic field configurations on the magneto-elastic oscillations of magnetars. For this purpose we have constructed magnetized equilibrium configurations generalizing the {\\small MAGSTAR} routine of the {\\small LORENE} library to include various descriptions of the current generating the magnetic field: {\\small MAGNETSTAR}\\footnote{ Our generalization of {\\small MAGSTAR} is publicly available in the standard {\\small LORENE} library.}. The oscillation spectrum, dominated by the continuum of the core, was studied by means of a semi-analytic model \\citep{Cerda2009,Gabler2012} and numerical GRMHD simulations.  Table\\,\\ref{tab_spindown} summarizes the main results for the magnetic field configurations considered in this paper. For each model we identify the fundamental Alfv\\'en QPO as the turning point in the Alfv\\'en continuum at the lowest frequency. We identify this mode with some of the observed QPOs, the $30\\,$Hz oscillation of SGR 1806-20 and $28\\,$Hz of SGR 1900+1, and we estimate the {\\it equivalent dipole magnetic field} strength of each configuration to match both frequencies. We define the {\\it equivalent dipole magnetic field}, $\\bar{B}$, as the magnetic field strength at the surface of a Newtonian, uniformly magnetized sphere having the same dipole magnetic moment\\footnote{The value of $m$ is provided in the output of the {\\small LORENE} library, and the corresponding magnetic field strength can be computed from $m=\\bar{B} R^3$. For the ring current $m$ is given by $m=IA$, where $I$ is the current and $A$ the area of the current loop.} $m$ as the configuration we want to describe, which can be directly compared with the magnetic field estimates from spin-down measurements, through $\\bar B[\\mathrm{G}]>3.2\\times10^{19}\\left(P[\\mathrm{s}]\\dot P\\right)^{1/2}$ \\citep{Lorimer2004}. Thus, the magnetic field estimates in Table\\,\\ref{tab_spindown} are directly comparable to estimates of the magnetic field from the spin down of the neutron star. \\begin{table} \\begin{center} \\begin{tabular}{c | c | c c} Model&QPO&$\\bar B_{15}^{28Hz}$&$\\bar B_{15}^{30Hz}$\\\\\\hline A0&U1&2.1&2.3\\\\ C$_{0.1}$&U1&1.3&1.4\\\\ C$_{10}$&U1&2.1&2.2\\\\ A1&U1&3.1&3.3\\\\ A1&U2&2.1&2.3\\\\ O&U1&3.6&3.8\\\\ O&U2&1.9&2.0\\\\ F&U1&2.4&2.5\\\\ \\hline R$_4$&U1&1.1&1.2\\\\ R$_4$&U2&0.85&0.93\\\\ R$_5$&U1&1.5&1.6\\\\ R$_5$&U2&0.80&0.85\\\\ R$_6$&U1&2.0&2.1\\\\ R$_7$&U1&2.7&2.8\\\\ R$_8$&U1&3.5&3.8\\\\\\hline $b_0=0$&U1&2.1&2.3\\\\ $b_0=1$&U1&2.2&2.4\\\\ $b_0=2$&U1&2.4&2.6\\\\ $b_0=5$&U1&2.9&3.1\\\\ $b_0=10$&U1&3.5&3.8\\\\ $b_0=20$&U1&3.7&4.0\\\\\\hline Q/D=0.1&U1&2.1&2.2\\\\ Q/D=0.5&U1&2.5&2.7\\\\ Q/D=0.5&U2&3.5&3.7\\\\ Q/D=0.5&U3&1.5&1.6\\\\ Q/D=1.0&U1$_N$&0.93&0.99\\\\ Q/D=1.0&U1$_S$&4.0&4.3\\\\ Q/D=2.0&U1$_N$&0.53&0.57\\\\ Q/D=2.0&U1$_S$&0.87&0.93\\\\ Q/D=5.0&U1$_N$&0.24&0.25\\\\ Q/D=5.0&U1$_S$&0.28&0.30\\\\ Q/D=10.0&U1$_N$&0.12&0.13\\\\ Q/D=10.0&U1$_S$&0.13&0.14\\\\ Q&U1&0.076&0.081\\\\\\hline \\end{tabular} \\caption{Equivalent field strength $\\bar B$ to match the frequencies at $f=28$ and $f=30\\,$Hz with the QPO indicated in the table. $\\bar B$ is defined by $m=\\bar B R^3$ for a uniformly magnetized sphere, where $m$ is the magnetic moment and $R$ is the radius of the star. The QPOs $Un_X$ are the Upper QPOs of the open field lines, $n$ indicates the number of the corresponding turning points labelled from $n=1$ at the polar axis to maximum $n=3$ at larger $\\theta$, and $X$ gives the division into northern ($X=N$) and southern ($X=S$) hemisphere, if there exist different QPOs in both hemispheres.}\\label{tab_spindown} \\end{center} \\end{table}    We summarize and discuss our main results next:   \\begin{itemize} \\item For poloidal, purely dipole-like configurations generated by a toroidal current inside the star (type II) we observe that the Alfv\\'en continuum always contains turning points at the polar axis (U1). Some models (A1, O, R$_4$, and R$_5$) have a second turning point (U2) in the region of open field lines inside the star. It seems very promising to identify some of the frequencies observed in the SGR giant flares, which are in the integer relation of 1:3:5 to each other\\footnote{$30$, $92$, and $150\\,$Hz for SGR 1806-20, and $28$, $84$, and $155\\,$Hz for SGR 1900+14}, with the fundamental Alfv\\'en Upper QPO at the polar axis and its second and fourth overtone. As can be seen in Table\\,\\ref{tab_spindown}, the required equivalent field strength is in the range of $0.8$ to $3.8\\times10^{15}\\,$G. This is consistent with the spin-down estimates of $\\bar B_{15}>2.4$ for SGR 1806-20 and $\\bar B_{15}>0.7$ for SGR 1900+14 (where $\\bar B_{15}$ is $\\bar B$ in units of $10^{15}$ G), respectively.  \\item The appearance of more than one turning point in the spectrum of some of the configurations allows us to relate additional observed frequencies to the torsional Alfv\\'en oscillations. Model O was constructed such that we can match all frequencies observed in SGR 1900+14. It is possible to identify the $28$\\,Hz with U1 at the polar axis, and the $84\\,$Hz$=3\\times28\\,$Hz with its second overtone. Additionally, the $53\\,$Hz QPO could be associated to U2 away from the polar axis at $\\chi\\sim5\\,$km (see blue line in Fig.\\,\\ref{fig_spectra}) and $155\\,$Hz $\\sim3\\times53\\,$Hz could be its second overtone. Interpreting the $30\\,$Hz QPO of SGR 1806-20 as U2, and the $92\\,$ and $150\\,$Hz ones as the second and fourth overtones, one could associate one of the lower frequencies, $18$ or $26\\,$Hz, with U1.  \\item We performed numerical simulations to test whether pure crustal shear oscillations can survive in the gaps of the Alfv\\'en continuum, as it was proposed by \\cite{vanHoven2011}. For this purpose we constructed a particular model (F) with a very flat spectrum which has very large gaps in between successive overtones. In none of our simulations we could find crustal shear modes. Instead, % predominantly Alfv\\'en oscillations of the continuum were present.  \\item For mixed poloidal-toroidal magnetic field configurations (type I) where the toroidal field is confined in the region of closed poloidal magnetic field lines inside the star, we estimate the part of the Alfv\\'en spectrum of axial oscillations in the region of open field lines, {where the toroidal component vanishes.} Qualitatively, the spectrum is similar to the one of purely poloidal fields (type II) in that region. With increasing field strength of the toroidal component, the poloidal field strength near the closed field lines increases, too. This also means that the surface field becomes stronger. Table\\,\\ref{tab_spindown} (3rd set) shows the block for the different configurations $b_0$ we see that the necessary  {\\it equivalent dipole magnetic field} strength to match the observed QPO frequencies in the spectrum increases with increasing toroidal field, because of the stronger increase of the poloidal field strength close to the surface, while the field in the center is only mildly affected.  \\item Configurations with an additional quadrupole-like component have a richer spectrum when this component is of comparable strength as the dipole-like one. For $Q/D=0.5$, we find three turning points in the spectrum. The additional turning points can produce oscillations which could be used to match a larger number of observed frequencies. For a quadrupole dominated configuration $Q/D\\gtrsim1.0$, we find asymmetric oscillations in both hemispheres of the neutron star. These provide a source for additional QPO frequencies which could be matched to observations. Models with $Q/D>>1$, i.e. dominantly quadrupole-like configurations, are currently not favoured, since they require magnetic field strengths lower than standard spin-down estimates. The presence of a large quadrupole-like component in magnetars, comparable to the poloidal component inside the star is not ruled out by our model. In fact the possible existence of this strong component could be related to the amplification of magnetic fields due to dynamo action, in initially rapidly rotating neutron stars in the standard magnetar formation scenario \\citep{Duncan1992}. Once the dynamo stops, the magnetic field evolves mainly due to Hall drift \\citep{Goldreich1992} on a characteristic timescale \\begin{equation} \\tau_{\\rm Hall} = 5\\,10^4 \\left( \\frac{L}{1 {\\rm km}}\\right) \\left( \\frac{10^{16} {\\rm G}}{B} \\right) yr, \\end{equation} where $B$ is the typical magnetic field strength in the magnetar interior, and $L$ is the typical length scale of the magnetic field loops. Under such conditions, a quadrupole-like component of similar strength as the dipole-like one could still be present in magnetars, if it was present at birth.  \\item Finally, we investigated magnetic field configurations limited to the crustal region. These configurations could be realized if the neutron star core is a superconductor of type I. In this case the shear modes cannot be damped into the core of the neutron star and they might survive sufficiently long to become observable. However, our simulation results suggest that these models provide no viable explanation for the observed frequencies in SGRs. We showed that the influence of the magnetic field strength on the frequencies of the magneto-elastic QPOs can be cast into a semi-empiric formula, which also contains the effect of different shear moduli. With this relation, we are able to scan a wide range of parameters of magnetic field strength and magnitude of the shear modulus. The two possible types of QPOs are not favoured because: i) On one hand, the structure and in particular the frequencies of predominantly shear modes in the zero-magnetic-field case change significantly with increasing magnetic field, such that the lowest observed frequencies of $f\\lesssim30Hz$ cannot be reached for the currently expected values of the shear modulus and the magnetic field strength. ii) On the other hand, predominantly Alfv\\'en oscillations will be quickly damped by phase mixing due to strong gradients in the spectra and an absence of turning points. For different EoS we obtain similar results. As expected, the core EoS does not play a large role in the determination of the critical magnetic field strength at which the frequencies begin to increase significantly with the magnetic field. The NV crust EoS leads to larger crusts with stronger shear moduli compared to the DH EoS. This, in turn, causes the necessary magnetic field to change the predominantly shear QPOs significantly to be up to $2\\times10^{14}\\,$G stronger than for the DH EoS. For quadrupole- or octupole-like dominated fields, the increase in frequencies of predominantly shear QPOs is less strong compared to dipole-like fields. However, from the neutron star spin-down measurement one has estimates of the dipole-like component of $B_{15}\\sim1$. At this dipole-like field strength our previous arguments already apply and the frequencies are shifted to $f\\gtrsim30\\,$Hz. \\end{itemize}  To summarize our findings, we have explored a large parameter space of possible magnetic field configurations of neutron stars. We discovered additional features, such as new turning points in the spectra for purely poloidal configurations, which have the potential to explain more of the observed QPO frequencies. Configurations with a toroidal component, which is confined to the closed poloidal field lines, have very similar spectra compared to those of purely poloidal fields. Furthermore, our results do not favour the assumption that the magnetic fields are confined to the crust.  Our magnetic field configurations do not include fields which have a toroidal component extending throughout the whole neutron star, which may be realized during a giant flare where strong toroidal fields are expected in the exterior \\citep{Beloborodov2009, Fernandez2007, Nobili2008}. For such configurations \\cite{Colaiuda2011b} find discrete Alfv\\'en oscillations. However, these fields require currents at the surface of the neutron star, which are difficult to model (and are set ad hoc in other studies). We postpone such considerations to future work.  The origin of the high frequency QPOs with $f=625\\,$Hz or higher remains open. In the magneto-elastic model the crustal $n=1$ shear modes are damped less efficiently than the $n=0$ modes, however, they still disappear on timescales much too short to be observable \\citep{Gabler2012, vanHoven2011}. One potential alternative solution of this problem is related to non-axisymmetric Alfv\\'en oscillation of superfluid stars \\citep{Passamonti2012}. Also oscillations of the magnetospheric field may play a role at these high frequencies.  In a previous work \\citep{Gabler2012} we have studied the effect of different EoS for one particular magnetic field configuration. The differences in the estimated magnetic field strength required to match observations that are caused by changing the EoS are of the same order (factor of a few) as the differences caused by assuming different magnetic field configurations. One of the main open questions in order to match observations is the effect of superfluidity in the core of the neutron star, which probably leads to lower (and thus probably more realistic) estimates of the magnetic field strength \\citep[see][for recent progress on this topic]{Passamonti2012}. Furthermore, a model of how the coupling of the interior oscillations to the magnetosphere can lead to a modulation of the X-ray emission needs to be developed. We plan to address these two points in future studies."
    },
    "1208/1208.3927_arXiv.txt": {
        "abstract": "The equations governing null and timelike geodesics are derived within the 3+1 formalism of general relativity. In addition to the particle's position, they encompass an evolution equation for the particle's energy leading to a 3+1 expression of the redshift factor for photons. An important application is the computation of images and spectra in spacetimes arising from numerical relativity, via the ray-tracing technique. This is illustrated here by images of numerically computed stationary neutron stars and dynamical neutron stars collapsing to a black hole. ",
        "introduction": "The computation of trajectories of photons or test-mass particles in the Kerr metric is a topic of major importance in relativistic astrophysics. This notably allows the investigation of spacetime properties around black holes (see e.g.~\\cite{fabian89,karas92,paumard08,hamaus09,dovciak04,StrauVAGP12} and references therein), the aim being to determine the black hole's mass and spin and to test general relativity (GR). Photons and test-mass particles follow spacetime null and timelike geodesics, respectively. Their motion is thus governed by the so-called \\emph{geodesic equation}.  However, within the framework of metric theories, strong tests of GR require to compare the Kerr geometry with geometries generated by alternative models of compact objects. The metric is then generally not known analytically and must be computed numerically. Rotating gravastars and boson stars are examples of such objects. Numerical metrics being almost exclusively computed using the 3+1 formalism of GR (see e.g. \\cite{Gourg12}), it is quite useful to derive the geodesic equation within this framework. This is a way to build an optimized ray-tracing algorithm.  In addition to the GR tests around astrophysical black holes, another field of application is the visualization of computer-generated spacetimes, resulting from numerical relativity studies of sources of gravitational radiation, such as gravitational collapse or coalescing binary compact objects \\cite{Alcub08,BaumgS10}. Such spacetimes are generally computed with the 3+1 formalism, and this motivates the design of a ray-tracing algorithm based on a 3+1 geodesic equation.  A ray-tracing code capable of using a 3+1 metric, \\texttt{GYOTO}, has recently been developed in our group \\cite{vincent11}. This code is written in C++, is open source and can be freely downloaded from \\cite{gyoto}. It computes null and timelike geodesics, both in the Kerr metric and in any numerically computed spacetime. The goal of this article is to derive the 3+1 geodesic equation that allows \\texttt{GYOTO} to compute images and spectra in numerically generated spacetimes, and to give the first examples of astrophysical interest of this capability. To our knowledge, in previous works, the geodesic equation has only been integrated in numerical spacetimes for the purpose of locating event horizons \\cite{HugheKWWST94,LibsoMSSW96,CohenPS09}, but not to form images nor to compute spectra.  The plan of the article is as follows. Section~\\ref{sec:31eq} derives the 3+1 geodesic equation. Section~\\ref{sec:red} derives the 3+1 expression of the redshift factor, useful for ray-tracing computations. Section~\\ref{sec:app} presents the first applications of ray-tracing in numerical spacetimes considering stationary and collapsing neutron star spacetimes. Finally, Section~\\ref{sec:conc} gives conclusions and perspectives for future works. ",
        "conclusions": "\\label{sec:conc}  We have re-expressed the geodesic equation within the framework of the 3+1 formalism of general relativity, obtaining equations (\\ref{e:dEdt}), (\\ref{e:systXV}) and (\\ref{e:redshift}). Equation~(\\ref{e:dEdt}), ruling the evolution of the particle energy with respect to Eulerian observers, has already been derived (in an equivalent form) by Merl{\\'\\i}n and Salgado \\cite{merlin11}. On the other hand, the system (\\ref{e:systXV}) for the position of the geodesic and the redshift formula (\\ref{e:redshift}) are novel. In particular, (\\ref{e:systXV}) significantly differs from previous 3+1 geodesic equations in  the literature \\cite{HugheKWWST94}, as discussed in Section~\\ref{s:3p1_geod_coord}. The 3+1 equations have been implemented in the ray-tracing code \\texttt{GYOTO} \\cite{vincent11,gyoto}, which has enabled us to compute images of stationary and collapsing neutron star numerical spacetimes generated by the \\texttt{LORENE/nrotstar} \\cite{gourgoulhon10b,lorene} and \\texttt{CoCoNuT}~\\cite{dimmelmeier05} codes.  Future work will be devoted to the development of ray-tracing computations in numerical spacetimes for astrophysically relevant problems. In particular, we shall try to carry on computations of images and spectra of astrophysical phenomena in the vicinity of compact objects, which can be alternative to black holes. This work is of particular interest in the perspective of a direct test of the nature of the central compact object of the Galaxy, Sgr~A*~(see the review \\cite{genzel10}, and in particular the section devoted to the alternatives to the black hole case).   The capability of \\texttt{GYOTO} to integrate geodesics in numerical spacetimes will be very interesting too, in order to visualize spacetimes, be it binary black holes spacetimes, binary neutron stars spacetimes, black hole -- neutron star binary spacetimes, or any other interesting metric (see reviews on these topics by~\\cite{hinder10,mcwilliams11,shibata11,faber12}) . \\texttt{GYOTO} could be used to image a background sky of stars, or a simple coordinate grid, putting in light the effect of strong gravity on these background objects.      \\ack The authors thank the anonymous referees for interesting comments that allowed to improve the quality of the article. This work was supported by grants from R\\'egion Ile-de-France and by the ANR grant 06-2-134423 \\emph{M\\'ethodes math\\'ematiques pour la relativit\\'e g\\'en\\'erale}.   \\appendix"
    },
    "1208/1208.6505_arXiv.txt": {
        "abstract": "The most studied way to explain the current accelerated expansion of the universe is to assume the existence of dark energy; a new component that fill the universe, does not clumps, currently dominates the evolution, and has a negative pressure. In this work I study an alternative model proposed by Lima et al. \\cite{lima96}, which does not need an exotic equation of state, but assumes instead the existence of gravitational particle creation. Because this model fits the supernova observations as well as the $\\Lambda$CDM model, I perform in this work a thorough study of this model considering an explicit spatial curvature. I found that in this scenario we can alleviate the cosmic coincidence problem, basically showing that these two components, dark matter and dark energy, are of the same nature, but they act at different scales. I also shown the inadequacy of some particle creation models, and also I study a previously propose new model that overcome these difficulties. ",
        "introduction": "Currently the observational evidence coming from supernovae studies \\cite{SNIa}, cosmic background radiation fluctuations \\cite{cmbr}, and baryon acoustic oscillations \\cite{bao}, set a strong case for a cosmological (concordance) universe model composed by nearly $70$ percent of a mysterious component called dark energy, responsible for the current accelerated expansion, nearly $25$ percent of dark matter, which populates the galaxy halos and a small percentage (around $4$) is composed by baryonic matter. The nature of these two dark component remains so far obscure \\cite{DEDM}. In the case of dark matter, we have a set of candidates to be probed with observations and detections in particle accelerators, however the case of dark energy is more elusive. This component not only has to fill the universe at the largest scale homogeneously, and also has the appropriate order of magnitude to be comparable to dark matter, but also has to have a negative pressure equation of state, something that was assumed first to explain inflation in the early universe, inspired by high energy theories of particle physics, but which seems to be awkward to appeal for, at these low energy scales. The alternative way to account for the cosmic acceleration, in the framework of the standard model, is to consider a modification of general relativity at large scales \\cite{alterGR}.  Some years ago, Prigogine and co-workers \\cite{prigogine} presented a very interesting cosmological model where matter creation takes place without spoiling adiabatic expansion. This is possible by adding a work term due to the change of the particle number density $n$. This work actually suggested for the first time, a way to incorporate the particle creation process in the context of cosmology in a self consistent way. In fact, the original claim made by Zeldovich \\cite{zeldov70}, that gravitational particle production can be described phenomenologically by a negative pressure, is realize here in a beautiful way. A covariant formulation of the model was presented for the first time in \\cite{CLW90}.  Originally, these considerations were applied in the context of; the steady-state cosmological model, the warm inflationary scenario and within the standard inflationary scenario, during the reheating phase. Based on this work, the studies of cosmological models with matter creation \\cite{lima96} were initiated, and rapidly recognized to be potentially important to explain dark energy \\cite{harko}. In particular, in Ref.\\cite{zsbp01} the authors established a model where dark energy can be mimicked by self-interactions of the dark matter substratum. Within the same framework, models of interacting dark energy and dark matter were proposed \\cite{interact}. Actually, in these studies is possible to have consistently, a universe where matter creation proceeds within an adiabatic evolution. More recently, Lima et al.,\\cite{lima08} have presented a study of a flat cosmological model where a transition from decelerated to accelerated phase exist. They explicitly show that previous models considered \\cite{lima99}, does not exhibit the transition, and study the observational constraints on the model parameters. Also in \\cite{Steigman:2008bc} the authors modified slightly the model, adding explicitly a baryonic contribution, which enable them to have a transition from decelerated to accelerated expansion always.  In this work I consider the non flat extension of this matter creation cosmological model. This is important because, even in the $\\Lambda$CDM concordance model, there is no clear evidence that $\\Omega_{k}$ is zero, due to the well known degeneracy between $w(z)$ and $\\Omega_{k}$ \\cite{curvature}. Furthermore, any alternative model to dark energy then must consider a non flat assumption as a prior. This model enable us to explain the current acceleration of the universe expansion through a fictitious pressure component coming from changes in the dark matter particle number, without any exotic contribution. As a bonus, this model enable us to explain easily the cosmic coincidence problem; it is not strange to have a similar contribution from these two (commonly differently regarded) components, because they are just two aspects of one and the same component; the dark matter.  In the next section I derive the equations of motion in the case of matter creation. Then, in section III I study the case for non flat models  and the observational consequences in the models already knew. After this, I discuss a new particle creation model that resembles the $\\Lambda$CDM model case, showing the transition from a decelerated to an accelerated expansion without enter in conflict with the behavior for large z. ",
        "conclusions": ""
    },
    "1208/1208.3116_arXiv.txt": {
        "abstract": "{In this contribution we discuss about numerical modeling of nebulae. In particular we emphasize on the dynamical evolution of an H~{\\sc ii} region and the chemical structure of a Photodissociation region. We do this using the Smoothed Particles Hydrodynamics code \\texttt{SEREN} and the recently developed astrochemistry code \\texttt{3D-PDR}, respectively. We show an example application by simulating a cometary globule using these two codes.} ",
        "introduction": "\\label{sec:1}  Nebulae (H~{\\sc ii} regions) are large regions consisting of ionised gas, particularly of hydrogen. The source of ionisation is usually a single or multiple massive stars emitting ultraviolet radiation with photons carrying more energy than the ionisation potential ($h\\nu>13.6{\\rm eV}$). A nebula is mainly structured by three different parts; the ionised region (number density of $n\\le200\\,{\\rm cm}^{-3}$ and temperature of $T\\simeq10^4\\,{\\rm K}$), the photodissociation region\\footnote{known also as ``Photon Dominated Region''}  (PDR; $200 \\le n \\le 10^5\\,{\\rm cm}^{-3}$, $10-20\\le T\\le 10^4\\,{\\rm K}$), and the dark molecular region ($n\\ge 10^5\\,{\\rm cm}^{-3}$, $T\\simeq10-20\\,{\\rm K}$). Of particular interest are PDRs; they are ubiquitously present in the interstellar medium (ISM) consisting of predominantly neutral gas and dust illuminated by FUV radiation ($6\\le h\\nu\\le 13.6{\\rm eV}$) and they occur in any region of the ISM that is dense and cold enough to remain neutral but has too low column density to prevent the penetration of FUV photons. Over the past few decades, effort has been made to study numerically the physics and chemistry of nebulae. Due to computational speed and memory capacity issues, detailed three-dimensional simulations are a reality since the past 5-7 years only. Even nowadays the numerical codes are divided in two main categories; codes that study the dynamical evolution and codes that study the chemical structure. An integrated code that treats simultaneously detailed dynamics, UV propagation and chemistry offering a realistic temperature and therefore pressure structure in H~{\\sc ii} regions is still lacking (although significant effort in this direction has been made by \\cite{Glov10, Clar12, Hawo12}). ",
        "conclusions": ""
    },
    "1208/1208.0673_arXiv.txt": {
        "abstract": "{Observations of molecular spectral lines provide information on the gas kinematics and chemistry of star-forming regions.} {We attempt to achieve a better understanding of the gas distribution and velocity field around the deeply embedded Class 0 protostar SMM 3 in the Orion B9 star-forming region.} {Using the APEX 12-m telescope, we mapped the line emission from the $J=2-1$ rotational transition of two CO isotopologues, $^{13}$CO and C$^{18}$O, over a $4\\arcmin \\times 4\\arcmin$ region around Orion B9/SMM 3.} {Both the $^{13}$CO and C$^{18}$O lines exhibit two well separated velocity components at about 1.3 and 8.7 km~s$^{-1}$. The emission of both CO isotopologues is more widely distributed than the submillimetre dust continuum emission as probed by LABOCA. The LABOCA 870-$\\mu$m peak position of SMM 3 is devoid of strong CO isotopologue emission, which is consistent with our earlier detection of strong CO depletion in the source. No signatures of a large-scale outflow were found towards SMM 3. The $^{13}$CO and C$^{18}$O emission seen at $\\sim1.3$ km~s$^{-1}$ is concentrated into a single clump-like feature at the eastern part of the map. The peak H$_2$ column density towards a C$^{18}$O maximum of the low-velocity component is estimated to be $\\sim10^{22}$ cm$^{-2}$. A velocity gradient was found across both the $^{13}$CO and C$^{18}$O maps. Interestingly, SMM 3 lies on the border of this velocity gradient.} {The $^{13}$CO and C$^{18}$O emission at $\\sim1.3$ km~s$^{-1}$ is likely to originate from the ``low-velocity part'' of Orion B. Our analysis suggests that it contains high density gas ($\\sim10^{22}$ H$_2$ molecules per cm$^2$), which conforms to our earlier detection of deuterated species at similarly low radial velocities. Higher-resolution observations would be needed to clarify the outflow activity of SMM 3. The sharp velocity gradient in the region might represent a shock front resulting from the feedback from the nearby expanding \\ion{H}{ii} region NGC 2024. The formation of SMM 3, and possibly of the other members of Orion B9, might have been triggered by this feedback.} ",
        "introduction": "The protostellar phase of low-mass star formation begins when a starless (prestellar) core collapses, and, after a hypothesised short-lived first-hydrostatic core stage (\\cite{larson1969}; \\cite{masunaga1998}), a stellar embryo forms in its centre (the so-called second hydrostatic core; e.g., \\cite{masunaga2000}). The dense cores harbouring the youngest protostars are known as the Class 0 objects (\\cite{andre1993}, 2000). In these objects, most of the system's mass resides in the dense envelope, i.e., $M_{\\rm env}\\gg M_{\\star}$, where $M_{\\star}$ is the mass of the central protostar. For this reason, Class 0 objects, or at least the youngest of them, are expected to still represent the initial physical conditions prevailing at the time of collapse phase. Class 0 objects are characterised by accretion-powered jets and mole\\-cular outflows, which can be very powerful and highly collimated (e.g., \\cite{bontemps1996}; \\cite{gueth1999}; \\cite{arce2005}, 2006; \\cite{lee2007}). The statistical lifetime of the Class 0 stage is estimated to be $\\sim1\\times10^5$ yr (\\cite{evans2009}; \\cite{enoch2009}), but the exact duration of this embedded phase of evolution can be highly dependent on the initial/environmental conditions (e.g., \\cite{vorobyov2010}).  The target source of the present study is the Class 0 protostellar core SMM 3 in the Orion B9 star-forming region, which was originally discovered by Miettinen et al. (2009; Paper I) through LABOCA 870-$\\mu$m dust continuum mapping of the region. SMM 3 is a strong submm emitting dust core ($S_{870}\\simeq2.5$ Jy) that is associated with a weak \\textit{Spitzer} 24-$\\mu$m point source ($S_{24}\\simeq5$ mJy), and a 3.6 Jy point source at 70 $\\mu$m. Using the Effelsberg 100-m telescope NH$_3$ observations, Miettinen et al. (2010; Paper II) derived the gas kinetic temperature of $T_{\\rm kin}=11.3\\pm0.8$ K in SMM 3. Using this temperature, the core mass was determined to be $7.8\\pm1.6$ M$_{\\sun}$, and its volume-averaged H$_2$ number density was estimated to be $1.1\\pm0.2\\times10^5$ cm$^{-3}$. In the SABOCA 350-$\\mu$m mapping of Orion B9 by Miettinen et al. (2012; Paper III), SMM 3 was found to be by far the strongest source in the mapped area ($S_{350}\\simeq5.4$ Jy). We also found that it contains two subfragments, or condensations (we called SMM 3b and 3c), lying about $36\\arcsec-51\\arcsec$ in projection from the central protostar. These correspond to 0.08--0.11 pc or $\\sim1.7-2.3\\times10^4$ AU at $d=450$ pc\\footnote{In this paper, we adopt a distance of 450 pc to the Orion giant molecular cloud (\\cite{genzel1989}). The actual distance may be somewhat smaller as, for example, Menten et al. (2007) determined a trigonometric parallax distance of $414\\pm7$ pc to the Orion Nebula.}. Because the thermal Jeans length of the core is $\\lambda_{\\rm J}=0.07$ pc, we suggested that the core fragmentation into condensations can be explained by thermal Jeans instability. Using the 350/870-$\\mu$m flux density ratio, we determined the dust temperature of the core to be $T_{\\rm dust}=10.8_{-2.6}^{+5.7}$ K, which is very close to $T_{\\rm kin}$ within the error bars. The revised spectral energy distribution (SED) of the core yielded a very low dust tempe\\-rature of 8 K, and a bolometric luminosity of $L_{\\rm bol}=1.2\\pm0.1$ L$_{\\sun}$. The latter is very close to the median luminosity of protostars in nearby star-forming regions, i.e., $L_{\\rm med}=1.5^{+0.7}_{-0.4}$ L$_{\\sun}$ (\\cite{enoch2009}; \\cite{offner2011}). In Paper III, we also studied the che\\-mistry of SMM 3. We derived a large CO depletion factor of $f_{\\rm D}({\\rm CO})=10.8\\pm2.2$, and a high level of deuterium fractionation, i.e., a ${\\rm N_2D^+}/{\\rm N_2H^+}$ column density ratio of $0.338\\pm0.092$. In Fig.~\\ref{figure:SMM3}, we show the LABOCA 870-$\\mu$m, SABOCA 350-$\\mu$m, and \\textit{Spitzer} 4.5/24-$\\mu$m images towards SMM 3. In Table~\\ref{table:SMM3}, we provide an overview of the physical and chemical properties of SMM 3 derived in our previous papers.  In this paper, we discuss the results of our $^{13}$CO and C$^{18}$O mapping observations of the environment of SMM 3. We ana\\-lyse the structure of the mapped region as traced by emission from the $J=2-1$ rotational transition of the above CO isotopologues. The rest of the present paper is organised as follows. Observations and data reduction are described in Sect.~2. Mapping results and analysis are presented in Sect.~3. In Sect.~4, we discuss our results, and in Sect.~5, we summarise and conclude the paper.  \\begin{figure}[!h] \\centering \\resizebox{0.8\\hsize}{!}{\\includegraphics{LABOCA.ps}} \\resizebox{0.8\\hsize}{!}{\\includegraphics{SABOCA.ps}} \\resizebox{0.8\\hsize}{!}{\\includegraphics{Spitzer.ps}} \\caption{LABOCA 870-$\\mu$m (\\textit{top}), SABOCA 350-$\\mu$m (\\textit{middle}), and a \\textit{Spitzer} IRAC/MIPS two-colour composite image (\\textit{bottom}; 4.5 $\\mu$m in green and 24 $\\mu$m in red) of the Class  0 protostellar core SMM 3 in Orion B9. The LABOCA and \\textit{Spitzer} images are shown with linear scaling, while the SABOCA image is shown with a square-root scaling to improve the contrast between bright and faint features. The LABOCA contours, plotted in white, go from 0.1 ($\\sim3.3\\sigma$) to 1.0 Jy~beam$^{-1}$ in steps of 0.1 Jy~beam$^{-1}$. The SABOCA contour levels, plotted in green, start at $3\\sigma$ and are 0.18 Jy~beam$^{-1}\\times [1,\\,2,\\,4,\\,6,\\,8,\\,10,\\,12,\\,14,\\,16]$. In the bottom panel, the first SABOCA contour, i.e., the $3\\sigma$ emission level, is plotted in white to guide the eye, and the white cross indicates the SABOCA peak position of SMM 3. The small subcondensations SMM 3b and 3c discovered in Paper III are labeled in the middle panel. The green plus sign shows the target position of our previous molecular-line observations (i.e., the submm peak position of the LABOCA map before adjusting the pointing; see Paper III for details). A scale bar indicating the 0.05 pc projected length is shown in the bottom left of the top panel, with the assumption of a 450 pc line-of-sight distance. The effective LABOCA and SABOCA beams, $\\sim20\\arcsec$ and $10\\farcs6$, are shown in the lower right corners of the corresponding panels.} \\label{figure:SMM3} \\end{figure}  \\begin{table} \\renewcommand{\\footnoterule}{} \\caption{Summary of the properties of SMM 3.} \\begin{minipage}{1\\columnwidth} \\centering \\label{table:SMM3} \\begin{tabular}{c c} \\hline\\hline Parameter & Value\\\\ \\hline {\\bf SMM 3} \\\\ $\\alpha_{2000.0}$\\tablefootmark{a} & 05$^{\\rm h}$ 42$^{\\rm m}$ $45\\fs3$ \\\\ $\\delta_{2000.0}$\\tablefootmark{a} & -01\\degr 16\\arcmin 16\\arcsec \\\\ ${\\rm v}_{\\rm LSR}$\\tablefootmark{b} & $8.68\\pm0.06$ km~s$^{-1}$\\\\ $R_{\\rm eff}$\\tablefootmark{c} & 13\\farcs3 (0.03 pc)\\\\ $T_{\\rm kin}$\\tablefootmark{d} & $11.3\\pm0.8$ K\\\\ $T_{\\rm dust}$\\tablefootmark{e} & $10.8_{-2.6}^{+5.7}$ K\\\\ $T_{\\rm dust}^{\\rm SED,\\, cold}$ & 8.0 K\\\\ $\\sigma_{\\rm NT}$\\tablefootmark{d} & $0.14\\pm0.003$ km~s$^{-1}$\\\\ $\\sigma_{\\rm NT}/c_{\\rm s}$\\tablefootmark{d} & $0.7\\pm0.04$ \\\\ $M$\\tablefootmark{f} & $7.8\\pm1.6$ M$_{\\sun}$/$2.1\\pm0.8$ M$_{\\sun}$\\\\ $\\alpha_{\\rm vir}$\\tablefootmark{g} & $0.5\\pm0.1$\\\\ $N({\\rm H_2})$\\tablefootmark{f} & $8.4\\pm1.1\\times10^{22}$/$1.0\\pm0.3\\times10^{23}$ cm$^{-2}$\\\\ $\\langle n({\\rm H_2}) \\rangle$\\tablefootmark{f} & $1.1\\pm0.2\\times10^5$/$4.0\\pm1.5\\times10^5$ cm$^{-3}$\\\\ $L_{\\rm bol}=L_{\\rm cold}+L_{\\rm warm}$ & $(0.3\\pm0.1)+(0.9\\pm0.1)=1.2\\pm0.1$ L$_{\\sun}$\\\\ $L_{\\rm submm}/L_{\\rm bol}$\\tablefootmark{h} & 0.1 \\\\ $f_{\\rm D}({\\rm CO})$ & $10.8\\pm2.2$ \\\\ $[{\\rm N_2D^+}]/[{\\rm N_2H^+}]$ & $0.338\\pm0.092$ \\\\ {\\bf SMM 3b} \\\\ $\\alpha_{2000.0}$\\tablefootmark{a} & 05$^{\\rm h}$ 42$^{\\rm m}$ $47\\fs6$ \\\\ $\\delta_{2000.0}$\\tablefootmark{a} & -01\\degr 16\\arcmin 24\\arcsec \\\\ $N({\\rm H_2})$\\tablefootmark{i} & $0.7\\pm0.2\\times10^{22}$ cm$^{-2}$\\\\ {\\bf SMM 3c} \\\\ $\\alpha_{2000.0}$\\tablefootmark{a} & 05$^{\\rm h}$ 42$^{\\rm m}$ $48\\fs6$ \\\\ $\\delta_{2000.0}$\\tablefootmark{a} & -01\\degr 16\\arcmin 32\\arcsec \\\\ $N({\\rm H_2})$\\tablefootmark{i} & $0.7\\pm0.2\\times10^{22}$ cm$^{-2}$\\\\ \\hline \\end{tabular} \\tablefoot{\\tablefoottext{a}{SABOCA 350-$\\mu$m peak position.}\\tablefoottext{b}{The LSR velocity derived from optically thin C$^{17}$O$(2-1)$ line.}\\tablefoottext{c}{Effective radius of the ``main'' core as determined from the SABOCA 350-$\\mu$m map.}\\tablefoottext{d}{Derived from NH$_3$ data. $\\sigma_{\\rm NT}$ and $c_{\\rm s}$ are, respectively, the one dimensional non-thermal velocity dispersion and the isothermal sound speed.}\\tablefoottext{e}{Computed from the 350-to-870 $\\mu$m flux density ratio.}\\tablefoottext{f}{The first value refers to the LABOCA 870-$\\mu$m core, and the second one to the ``main'' core detected at 350 $\\mu$m.}\\tablefoottext{g}{Virial parameter defined by $\\alpha_{\\rm vir}=M_{\\rm vir}/M$.}\\tablefoottext{h}{$L_{\\rm submm}$ is the submm luminosity derived by integrating the SED longward of 350 $\\mu$m.}\\tablefoottext{i}{Calculated by making the assumption that $T_{\\rm dust}$ equals the $T_{\\rm kin}$ derived for the ``main'' core.}} \\end{minipage} \\end{table} ",
        "conclusions": "A $4\\arcmin \\times 4\\arcmin$ region around the Class 0 protostar SMM 3 in Orion B9 was mapped in $^{13}$CO and C$^{18}$O $J=2-1$ lines with the APEX 12-m telescope. Our main results and conclusions can be summarised as follows:  \\begin{enumerate} \\item Both lines exhibit two well separated velocity components: one at $\\sim1.3$ km~s$^{-1}$ and the other at $\\sim8.7$ km~s$^{-1}$. The latter is near the systemic velocity of SMM 3. The low-velocity component was already recognised in our previous studies, and it is believed to be related to the low-velocity part of Orion B. \\item The $^{13}$CO and C$^{18}$O emission are relatively widely distributed compared to the dust continuum emission traced by LABOCA. The LABOCA 870-$\\mu$m peak position of SMM 3 is not coincident with any strong $^{13}$CO or C$^{18}$O emission, which is in accordance with the high CO depletion factor derived earlier by us from C$^{17}$O$(2-1)$ ($f_{\\rm D}\\simeq10.8$). The CO depletion factor derived from C$^{18}$O data is within a factor of two from the previous estimate, i.e., $f_{\\rm D}\\simeq5.6$. No evi\\-dence for a large-scale outflow activity, i.e., high velocity line wings, was found towards SMM 3. \\item The lower-velocity ($\\sim1.3$ km~s$^{-1}$) $^{13}$CO and C$^{18}$O emission are concentrated into a clump-like feature at the eastern part of the map. We estimate that the H$_2$ column density towards its C$^{18}$O maximum is $\\sim10^{22}$ cm$^{-2}$. Therefore, the lower-velocity gas seen along the line of sight is of high density, which is consistent with our earlier detection of, e.g., deute\\-rated molecular species (DCO$^+$, N$_2$D$^+$). \\item We observe a velocity gradient across the $^{13}$CO and C$^{18}$O maps along the NW-SE direction (some hint of that is also visible in the lower-velocity line maps). Interestingly, SMM 3 is projected almost exactly on the border of the velo\\-city jump. The sharp velocity-gradient border provides a strong indication that it represents an interaction zone of flow motions. \\item We suggest a possible scenario in which the formation of SMM 3, and likely some of the other dense cores in Orion B9, was triggered by an expanding \\ion{H}{ii} region of NGC 2024. This collect-and-collapse -type process might have been taken place some several times $10^5$ yr ago. The NGC 2024 region is known to be a potential site of induced, sequential star formation (e.g., \\cite{fukuda2000}, and references therein). The case of Orion B9 suggests that we may be witnessing the most recent event of self-propagating star formation around NGC 2024. Larger-scale molecular-line maps would be needed for a better understanding of the larger-scale velocity structure of the region. \\end{enumerate}"
    },
    "1208/1208.0390_arXiv.txt": {
        "abstract": "We study the condensation phenomenon for a system of charged bosons in the presence of an external magnetic field. We show that condensation happens for a definite critical temperature instead of through a diffuse phase transition. The essential ingredient, overlooked in previous analyses and accounted for in this work, is the treatment of the plasma screening effects by means of resummation. We compute the critical temperature, for the case in which the condensate is made of charged pions and for typical densities found in compact astrophysical objects, for small and large values of the magnetic field. We show that the magnetic field catalyzes the onset of condensation at very small and at large values of the magnetic field, and that for intermediate values the critical temperature for condensation is lower than for the zero magnetic field case. ",
        "introduction": "\\label{I}  The possibility that a charged pion condensate may occur in the interior of neutron stars has been repeatedly examined in the past. This possibility is raised by the large isospin imbalance between neutrons and protons which favors reactions that make neutrons decay into negative pions under appropriate conditions. The equilibrium thermodynamic conditions obeyed by a pion condensed state in dense neutron matter and in neutron stars have been discussed long ago. In particular, Ref.~\\cite{Baym} studies the criteria for the appearance of pion condensation in neutron matter in terms of the pion Green's function (for a general review on the physics of neutron stars see Ref.~\\cite{Glendenning}). The occurrence of a charged boson condensed phase without magnetic fields has also been extensively discussed in the literature. Refs.~\\cite{Sonandothers, Kogut, Splittdorf, Herpay:2008uw} study in-medium processes introducing an isospin chemical potential $\\mu _{I}$ at zero temperature in both phases ($ |\\mu _{I} | \\gtrless m_{\\pi }$, where $m_\\pi$ is the pion mass), analyzing the formation of a charged pion condensed phase. This phenomenon was discussed in electrically neutral dense quark matter in Refs.~\\cite{Ebert,Abuki,Andersen}. Finite temperature corrections, in the frame of chiral perturbation theory, have been considered in Ref.~\\cite{Villavicencio1}, extending the discussion also to other condensates, like the chiral condensate or the axial-isospin charge density condensate, in Ref.~\\cite{Villavicencio2}.   The situation becomes even more interesting when considering that neutron stars possess large magnetic fields whose effects should also be included when studying the condensation conditions. Recently, in Ref.~\\cite{Son2} the role played by the coupling of $\\pi^{0}$ to a magnetic field via the triangle anomaly has been considered, showing the emergence of a  $\\pi ^{0}$ domain wall for values of the magnetic field strength $B$ larger than a certain critical value. This could also happen for $(\\eta, \\eta ')$ states when $ B \\sim 10^{7} - 10^{19}$ G.  Magnetic fields can also play an important role in the dynamics of systems where charged pions are copiously produced, such as relativistic heavy-ion (RHI) collisions. Recently, the importance of large magnetic fields for the evolution of QCD matter produced in noncentral RHI collisions has been discussed in Ref.~\\cite{kharzeev} as well as their influence on the phase structure of QCD, with emphasis on the chiral symmetry restoration and deconfining transitions. In Ref.~\\cite{Mizher}, a discussion of the effective potential in the framework of the linear sigma model, coupled to quarks and/or Polyakov loop, suggests a richer structure of the strong interactions like, for example, a possible splitting between chiral symmetry restoration and deconfinement in the presence of magnetic fields. The influence of the external magnetic field on the formation of CP-odd domains in RHI collisions has also been discussed in Ref.~\\cite{Mizher2}. A decrease in the confining critical temperature was found in \\cite{Fraga:2012fs}, where a hadron-quark transition was studied within the MIT bag model. With the above ingredients put together, the theoretical study of a charged boson condensate with a finite chemical potential in the presence of magnetic fields becomes even more relevant. Although this is an old problem, the results from several approaches vary in their conclusions. For instance, it was long ago argued that a nonrelativistic Bose-Einstein gas of charged particles does not condense in the presence of a magnetic field, regardless of how weak the field may be \\cite{Schafroth}. This result motivated the search for conditions where condensation could take place with magnetic fields, in particular to study whether this could happen extending the number of spatial dimensions \\cite{May, Daicic, Elmfors}. In the nonrelativistic case, treating the dimensionality $d$ of the system as a continuous variable, it was shown in Ref.~\\cite{May} that condensation can happen only for $d>4$. For pairs of bosons or fermions and in the relativistic case, it was shown in Ref.~\\cite{Daicic} that for the case when $d$ is taken as an integer, condensation happens for odd $d\\geq 5$. A similar conclusion was reached in Ref.~\\cite{Elmfors}, although these authors also realized that the lowest Landau level can play the role of the ground state to accommodate a large charge density in the $d=3$ case.  The common feature of all of the above-mentioned analyses is the definition of the condensation condition which is taken as the equality of the chemical potential and the ground state energy. However, in the presence of a magnetic field, this condition leads to a divergence of the particle density for that state. Indeed, since for a constant magnetic field the energy levels separate into transverse and longitudinal (with respect to the magnetic field direction) and the former are described in terms of discrete energy levels, the divergence of the Bose-Einstein distribution when the chemical potential is equal to the lowest energy level can only be cured in a larger than $d=4$ number of spatial dimensions.  The implications of this condition were recognized in Ref.~\\cite{Perez} where it was argued that when the temperature $T$ is much lower than $eB$ one can already consider that the system occupies only the lowest Landau level. This means that the value for the chemical potential to compute the ground state density does not need to be equal to the lowest energy. In this picture the occupation of this state occurs without the need of having a critical temperature, that is, the system undergoes a {\\it diffuse phase transition}.  Nevertheless one can argue that if in the absence of a magnetic field the system is already in the condensed phase with a macroscopic fraction of the population occupying the lowest energy level, a slow turning on of the magnetic field should not lead to the instantaneous destruction of the condensate. Put in equivalent terms, the onset of condensation for small magnetic fields should be a phenomenon that takes place at a given critical temperature $T_c$ since it does so in the limit of a vanishing magnetic field and the presence of a small one cannot drastically change the picture. To implement this idea, one should keep in mind that the chemical potential is not a number that can arbitrarily be set to take a specific value but rather, a function of the thermodynamic variables such as temperature and density. Its value should be determined by demanding that the ground state is populated by a finite charge density. The missing ingredient that bridges the gap in the analysis is to consider the plasma screening effects, which are of course needed since we are dealing with infrared phenomena where {\\it the effective mass} is small or may even vanish.  In this work we study the conditions for the onset of a condensed phase for a charged boson system, in the presence of an external magnetic field. To mimic the situation where there is an isospin imbalance,  we introduce a finite  chemical potential $\\mu$. For the description, we resort to model the boson system in terms of a theory of a charged scalar with quartic self-interactions. We show that for small and large values of the magnetic field, the system presents the magnetic catalysis phenomenon~\\cite{catalysis}; that is, that the formation of the condensate is favored by the presence of the magnetic field. This phenomenon has also been found in the context of the Nambu-Jona-Lasino model at $T=0$~\\cite{Klevansky2} and in (2+1) dimensions both at $T=0$ and $T\\neq 0$~\\cite{Klimenko}, where it was shown that even the presence of an arbitrary small magnetic field breaks the chiral invariance of the models. A main result of our work is to show that when including the plasma screening effects, there is a well-defined critical temperature associated with the onset of condensation. A similar calculation, using optimized perturbation theory, albeit without the introduction of a chemical potential, was done in Ref.~\\cite{duarte}. The authors found that the phase transition is always second order and the magnetic catalysis phenomenon is present for all values of the magnetic field. They also found that the critical temperature increases with increasing values of the magnetic field.  The work is organized as follows: In Sec.~\\ref{II} we find the lowest energy state where condensation happens and define the order parameter for the transition. In Sec.~\\ref{III} we compute the one-loop corrections to the grand potential and set up the discussion for the onset of the condensation phenomenon in terms of the existence of a large but finite charge density in the ground state. In Sec.~\\ref{IV} we revisit the description of the onset of condensation when corrections from interactions are accounted for. We take the limit $B\\to 0$ and point out the need to include plasma screening effects by means of resummation, even in this case. In Sec~\\ref{V} we explicitly compute the resummed self-energy for finite $B$ in the low temperature approximation in the limits of small and large magnetic fields. This self-energy is then used in Sec.~\\ref{VI} to compute the critical temperature for condensation when the charged bosons are taken as pions, for typical densities in compact astrophysical objects such as neutron stars. We finally summarize and conclude in Sec.~\\ref{VII}. ",
        "conclusions": "\\label{VII}  In this work we have studied the condensation phenomenon for a system of charged bosons in the presence of an external magnetic field. Contrary to what is commonly believed, we have shown that condensation happens at a definite critical temperature. The missing ingredient, overlooked in previous analysis and accounted for in this work, is the treatment of the plasma screening effects by means of resummation. We have explicitly computed the critical temperature, for typical densities found in compact astrophysical objects, for small and large values of the magnetic fields. We have shown that the magnetic field catalyzes the onset of the condensation at very small and large values of the magnetic field, agreeing in these regions with the case of zero chemical potential \\cite{duarte}. For intermediate values of the magnetic field, the critical temperature for condensation turns out to be lower than in the $B=0$ case.   Recall that although the term {\\it magnetic catalysis} usually refers to an enhancement of dynamical symmetry breaking by an external magnetic field, the phenomenon seems to be universal, appearing in different physical scenarios~\\cite{Shovkovy} such as the one treated in this work, namely, a condensate of scalar particles, as opposed to the more usually studied case of a fermion condensate.  We should also mention that the problem has been studied by lattice methods as well. The first lattice simulation for deconfinement and chiral symmetry restoration for two-flavor QCD, in the presence of a magnetic background, was done in Refs. \\cite{D'Elia:2010nq,D'Elia:2011zu}. The result was that the transition temperature significantly decreased with increasing magnetic field. A similar conclusion was presented in \\cite{Ilgenfritz}. In that work it was found that in the chirally broken phase, the chiral condensate increased monotonically with a growing magnetic field strength. In fact, in the chiral limit this behavior started linearly. In the same limit and in the chirally restored phase, the condensate vanished independent of the strength of the magnetic field. On the other hand, in Refs.~\\cite{balietal,Bali:2012zg} the effect of an external magnetic field on the finite temperature transition of QCD was considered. Thermodynamic observables including the chiral condensate and the susceptibility were measured. The result was that the transition temperature significantly decreased with increasing magnetic field. Such discrepancies can be originated by the fact that in Refs. \\cite{D'Elia:2010nq,D'Elia:2011zu,Ilgenfritz} the pion had a large mass. In the case of Refs. \\cite{balietal,Bali:2012zg} light pions as well as an improved lattice were used. These seemingly contrasting results call for a closer look at the phenomenon, in particular for the case of intermediate values of the magnetic field strength. This is work that we are currently pursuing and will be reported elsewhere.      \\bigskip \\noindent"
    },
    "1208/1208.1675_arXiv.txt": {
        "abstract": "{ The \\pao is a facility built to detect air showers produced by cosmic rays above $10^{17}$~eV.  During clear nights with a low illuminated moon fraction, the UV fluorescence light produced by air showers is recorded by optical telescopes at the Observatory.  To correct the observations for variations in atmospheric conditions, atmospheric monitoring is performed at regular intervals ranging from several minutes (for cloud identification) to several hours (for aerosol conditions) to several days (for vertical profiles of temperature, pressure, and humidity).  In 2009, the monitoring program was upgraded to allow for additional targeted measurements of atmospheric conditions shortly after the detection of air showers of special interest, e.\\,g., showers produced by very high-energy cosmic rays or showers with atypical longitudinal profiles.  The former events are of particular importance for the determination of the energy scale of the Observatory, and the latter are characteristic of unusual air shower physics or exotic primary particle types.  The purpose of targeted (or ``rapid'') monitoring is to improve the resolution of the atmospheric measurements for such events.  In this paper, we report on the implementation of the rapid monitoring program and its current status.  The rapid monitoring data have been analyzed and applied to the reconstruction of air showers of high interest, and indicate that the air fluorescence measurements affected by clouds and aerosols are effectively corrected using measurements from the regular atmospheric monitoring program.  We find that the rapid monitoring program has potential for supporting dedicated physics analyses beyond the standard event reconstruction. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.4861_arXiv.txt": {
        "abstract": "We present the first results in the search for relativistic magnetic monopoles with the IceCube detector, a subsurface neutrino telescope located in the South Polar ice cap containing a volume of 1 km$^{3}$. This analysis searches data taken on the partially completed detector during 2007 when roughly 0.2 km$^{3}$ of ice was instrumented.  The lack of candidate events leads to an upper limit on the flux of relativistic magnetic monopoles of $\\Phi_{\\mathrm{90\\%C.L.}}\\sim 3\\e{-18}\\fluxunits$ for $\\beta\\geq0.8$. This is a factor of 4 improvement over the previous best experimental flux limits up to a Lorentz boost $\\gamma$ below $10^{7}$. This result is then interpreted for a wide range of mass and kinetic energy values. ",
        "introduction": "Introduction} Magnetic monopoles are an important element in a complete picture of our universe.  Their existence would explain the quantization of electric (and magnetic) charge via the Dirac quantization equation $g=Ne/2\\alpha$ \\cite{Dirac}.  They appear as topological defects from symmetry breaking in Grand Unified Theories (GUTs) \\cite{tHooftPoly} with masses $\\sim10^{4}-10^{17}\\,$GeV \\cite{Wick}, depending on the breaking scheme. Additionally, they would bring a complete symmetry to Maxwell's equations.  Magnetic monopoles produced in the early universe via GUT symmetry breaking would be topologically stable and accelerated along magnetic field lines.  The universe is full of long range magnetic fields that would act upon the monopoles over their lifetime, likely imparting energies $\\sim10^{15}\\,$GeV \\cite{Wick}.  Therefore, magnetic monopoles below this energy scale should reach and travel through the Earth at relativistic speeds. A relativistic magnetic monopole moving through a transparent medium would produce copious amounts of Cherenkov light, $\\sim$8300 times a single muon in ice \\cite{Tom}. Thus, large Cherenkov detectors like IceCube are an ideal experiment to search for these particles.  The current best limits on the flux of magnetic monopoles at the 90\\% confidence level (C.L.) for relativistic speeds between $\\beta=0.8$ and Lorentz boost $\\gamma=10^{7}$ are set by the ANTARES detector \\cite{ANTARES} at the $\\sim 10^{-17}\\fluxunits$ scale.  This recent result is the first in this velocity range to surpass the results from the AMANDA detector \\cite{AMANDA}, IceCube's proof of concept, which set flux limits $\\sim3\\e{-17}\\fluxunits$.  ANTARES also searched for magnetic monopoles below the Cherenkov threshold but still energetic enough to knock off electrons that produce Cherenkov light.  This extension sets flux limits at the $\\sim5\\e{-17}\\fluxunits$ scale down to a speed of $\\beta=0.625$. For lower speeds, MACRO provides comprehensive flux limits $\\sim10^{-16}\\fluxunits$ \\cite{MACRO} for speeds down to $\\beta\\sim10^{-4}$ while flux limits at ultra-relativistic speeds are set by radio detectors RICE \\cite{RICE} and ANITA \\cite{ANITA} at the $\\sim10^{-19}\\fluxunits$ scale.  These are important as they are flux limits below the 'Parker Bound' \\cite{PB} ($\\sim 10^{-15}\\fluxunits$), an astrophysical flux limit derived by considering the survival of the galactic magnetic field in the presence of magnetic monopoles. More sophisticated calculations that consider velocity \\cite{Turner} relax the bound on relativistic magnetic monopoles above a mass of $10^{11}\\,$GeV due to the shortened time spent in the galactic field. However, an 'Extended Parker Bound' found by considering the survival of a modeled seed field still produces flux limits well below experiments, with $\\Phi\\sim10^{-16}(\\mathrm{Mass_{MP}})/(10^{17}\\,\\mathrm{GeV})\\,\\fluxunits$ \\cite{EPB}.  This paper describes the search for relativistic magnetic monopoles in data taken with the IceCube detector between May 2007 and April 2008. The analysis is optimized for magnetic monopoles with modest Lorentz boosts ($\\gamma\\le10$) and charge $g=1$. The derived flux limits are conservative upper bounds for magnetic monopoles with larger $\\gamma$ or charge, as these cases produce more light in the ice. The paper is organized as follows.  Section \\ref{sec:Det} describes the IceCube detector.  Section \\ref{sec:Sim} describes the simulation of background and signal. Section \\ref{sec:EvSel} defines the variables and outlines the steps used to discriminate signal events from background. Section \\ref{sec:Unc} summarizes the uncertainties. Section \\ref{sec:Res} presents the results for an isotropic flux of magnetic monopoles at the detector.  Section \\ref{sec:Dis} extends this result to an isotropic flux at the Earth's surface by considering the energy loss of magnetic monopoles through the Earth. This results in a final limit plot that is presented over a large range of magnetic mass and kinetic energy values.  This allows the result to remain agnostic towards the particular origin and energy gaining mechanism a magnetic monopole may possess.  Concluding remarks are presented in Section \\ref{sec:Conc}. ",
        "conclusions": ""
    },
    "1208/1208.3736_arXiv.txt": {
        "abstract": "Accounting for $\\sim$20\\% of the total QSO population, Broad Absorption Line QSOs are still an unsolved problem in the AGN context. They present wide troughs in the UV spectrum, due to material with velocities up to 0.2 c toward the observer. The two models proposed in literature try to explain them as a particular phase of the evolution of QSOs or as normal QSOs, but seen from a particular line of sight.\\\\ We built a statistically complete sample of Radio-Loud BAL QSOs, and carried out an observing campaign to piece together the whole spectrum in the cm wavelength domain, and highlight all the possible differences with respect to a comparison sample of Radio-Loud non-BAL QSOs. VLBI observations at high angular resolution have been performed, to study the pc-scale morphology of these objects. Finally, we tried to detect a possible dust component with observations at mm-wavelengths.\\\\ Results do not seem to indicate a young age for all BAL QSOs. Instead a variety of orientations and morphologies have been found, constraining the outflows foreseen by the orientation model to have different possible angles with respect to the jet axis. ",
        "introduction": "\\vspace{0.3cm} Broad Absorption Line (BAL) QSOs are a particular and not yet understood class of AGN: their spectra show wide troughs towards the blue wing of some UV emission lines (\\mgii, \\aliii, \\siiv, \\civ), due to ionized gas with outflow velocities up to 0.2 c. They account for $\\sim$20\\% of the QSO population. Two main scenarios have been considered to explain their nature: the \\emph{orientation} model and the \\emph{evolutionary} one. The former, proposed by \\cite{Elvis}, foresees the BAL outflows to be present in all QSOs, but visible only when intercepting the line of sight of the observer, thus a particular orientation can be supposed for these objects. The latter model (\\cite{Briggs, Lipari}) explains the BAL phenomenon as a particular phase in the evolutionary sequence of QSOs, during which a dusty shell is being expelled from the AGN. In this case a greater amount of dust emission should be present with respect to \\emph{normal} QSOs. A third model from Punsly (\\cite{Punsly1,Punsly2}) proposes that the BAL phenomenon is due to polar winds above the inner jet, thus would be visible only for lines of sight close to the jet axis.  In this framework the radio emission is an additional diagnostic tool to test the proposed models. We selected a sample of Radio-Loud BAL QSOs, cross-correlating the fourth edition of the SDSS Quasar Catalogue (\\cite{Schneider}) with the FIRST catalogue (\\cite{Becker}), with a constraint of $S_{1.4~\\rm{GHz}}>30$ mJy for the flux density at 1.4 GHz and $1.7<z<4.7$ for the redshift, in order to have sufficient flux density for VLBI studies and to allow the classification as BAL QSO through the CIV and MgII lines in the optical band. The whole sample of 25 objects, together with a comparison sample of 34 non-BAL QSOs, is presented in \\cite{Bruni}. ",
        "conclusions": "\\vspace{0.5cm}  We can summarize our findings as follows: \\begin{itemize} \\item A variety of orientation for BAL QSOs have been found, both from spectral index and from VLBI morphology studies. \\item Radio spectra do not seem to indicate a young age, since the fraction of GPS sources is similar to non-BAL QSOs. Moreover, in some cases a low-frequency emission is present, suggesting an old radio component. \\item VLBI maps for 11 BAL QSOs show different morphologies and projected linear sizes under 1 kpc. Bigger linear sizes are not excluded for BAL QSOs, since 4 sources were already resolved with the VLA, resulting in projected linear sizes between 20 and 200 kpc. \\item Radio-Loud BAL QSOs do not seem to be dust-rich, since preliminary results only show $\\sim7\\%$ of the sources with a flux density excess in the mm-band, comparable with previous findings for the QSO population. \\end{itemize}  Excluding a young age for Radio-Loud BAL QSOs, the most probable scenario seems to be the one proposed by \\cite{Elvis}, but only if it can account for outflows with a variety of angles with respect to the jet axis, able to justify the different orientations found in radio observations.   \\subsection*"
    },
    "1208/1208.1505_arXiv.txt": {
        "abstract": "Based on sensitive CO measurements from HERACLES and \\hi\\ data from THINGS, we show that the azimuthally averaged radial distribution of the neutral gas surface density (\\sighi + \\sightwo) in 33 nearby spiral galaxies exhibits a well-constrained universal exponential distribution beyond $0.2\\times\\rtf$ (inside of which the scatter is large) with less than a factor of two scatter out to two optical radii \\rtf. Scaling the radius to \\rtf\\ and the total gas surface density to the surface density at the transition radius, i.e., where \\sighi\\ and \\sightwo\\ are equal, as well as removing galaxies that are interacting with their environment, yields a tightly constrained exponential fit with average scale length 0.61$\\pm$0.06\\,\\rtf. In this case, the scatter reduces to less than 40\\% across the optical disks (and remains below a factor of two at larger radii). We show that the tight exponential distribution of neutral gas implies that the total neutral gas mass of nearby disk galaxies depends primarily on the size of the {\\it stellar} disk (influenced to some degree by the great variability of \\sightwo\\ inside $0.2\\times\\rtf$). The derived prescription predicts the total gas mass in our sub-sample of 17 non-interacting disk galaxies to within a factor of two. Given the short timescale over which star formation depletes the \\htwo\\ content of these galaxies and the large range of \\rtf\\ in our sample, there appears to be some mechanism leading to these largely self-similar radial gas distributions in nearby disk galaxies. ",
        "introduction": "\\label{intro} Scaling relations are important because they are often thought to be manifestations of some underlying fundamental physical processes. They also make it possible to determine properties of a system with incomplete measurements to within the scatter of the scaling. The total neutral gas content of galaxies has long defied such characterizations because of the wide diversity of previously observed gas surface density distributions \\citep[see, e.g.,][]{helfer03}. This was at least partly due to the lack of comprehensive, resolved, extended and deep observations of atomic (\\hi) and molecular gas (\\htwo) across the disks of a significant sample of galaxies. This limitation has been overcome by the availability of new radio surveys providing sensitive \\hi\\ and CO (the standard tracer for \\htwo) radio data, tracing the gas content of many nearby galaxies out to large galactocentric radii \\citep[THINGS, HERACLES,][]{walter08,leroy09}.  With these data sets in hand, we can now re-assess the azimuthally averaged gas distributions (radial profiles) in nearby spiral galaxies. Both HI and CO profiles tend to show quite distinct behavior: the HI surface density distributions are often depressed in the centers, increasing in the inner disk and slowly declining in the outer regions.  Furthermore, the surface density in the outer regions can be affected by environmental factors such as interaction with companions or truncation due to the ram pressure stripping by the hot gas in galaxy clusters \\citep{cayatte94}. As for the molecular gas, it tends to fall off much more sharply than the HI in an approximately exponential fashion that varies from galaxy to galaxy \\citep{young82,regan01,leroy08,bigiel08}.  We show in this paper, however, that when the totality of the neutral (i.e., atomic and molecular) gas in normal spirals is considered based on the data from these new surveys, and the profiles are properly scaled, the gas exhibits a radial profile shape that is remarkably constant from galaxy to galaxy. The gas surface density is found to vary by no more than about a factor of two from this average profile across the optical disk and even beyond. We show that the scatter reduces even further when we eliminate galaxies that are interacting with their environment. ",
        "conclusions": "\\label{discussion}  Figure \\ref{fig3} shows that galaxies not interacting with their environment exhibit a tight exponential distribution of their total gas content when the radius is scaled to the optical radius and the surface density of the gas is scaled to the surface density of the gas at the transition radius. The scaling relation for disk galaxies we derive is:  \\begin{equation} \\frac{\\siggas}{\\Sigma_{\\rm trans}} = 2.1\\times e^{-1.65\\times r/r_{25}}, \\label{hih2} \\end{equation}  \\noindent where {\\mbox{$\\Sigma_{\\rm trans}$}} is the surface density of the gas at the transition radius. The factors $2.1$ and $-1.65$ are obtained from the y-intercept and slope of the exponential fit in Figure 3. This implies that for these disk galaxies the total mass of neutral gas, ${\\rm M_{gas}}$, is given by  \\begin{equation} {\\rm M_{gas}} = 2\\pi\\times 2.1\\times \\Sigma_{\\rm trans} \\times \\rtf^2 \\times {\\rm X}, \\label{mgas} \\end{equation}  \\noindent where the factor X depends on how far out from the center one integrates Equation \\ref{hih2}. Integrating out to the optical radius \\rtf\\ yields ${\\rm X} = 0.18$, out to 2$\\times$\\rtf\\ yields ${\\rm X} = 0.31$ and integrating to infinity yields ${\\rm X} = 0.37$.  Equation \\ref{mgas} shows that for any given maximum radius, ${\\rm M_{gas}}$ depends only on ${\\rtf}^2$ and ${\\Sigma_{\\rm trans}}$. ${\\Sigma_{\\rm trans}}$ does not vary a great deal from galaxy to galaxy and has a typical value of about 14\\,M$_\\odot$~pc$^{-2}$ \\citep[also compare][]{leroy08,bigiel08}. Thus, ${\\rm M_{gas}}$ depends primarily on ${\\rtf}^2$, which varies by two orders of magnitude for the galaxies plotted.  This leads to the surprising result that the mass of neutral hydrogen gas depends mostly on the size of the {\\it stellar} disk and that the gas arranges itself somehow into a distribution that is self-similar among galaxies. The result also implies that except for the region of a galaxy at $\\lesssim0.2\\,\\rtf$, which is highly variable, the total neutral gas mass of a disk galaxy at $z = 0$ can be estimated if \\rtf, i.e., the extent of the stellar disk, is known.  The variable inner regions of galaxies, are, however, quite small, and despite either depressions or large excesses of the molecular gas in these regions, they typically only account for a small fraction of the total neutral gas mass ($\\sim15\\%$ on average in our sample).  \\begin{figure*}[!t] \\epsscale{0.6} \\plotone{fig6.eps} \\caption{Total (\\hi+\\htwo) gas mass obtained from the average radial profile fit (Figure \\ref{fig3} and Equation \\ref{mgas}) versus the mass obtained directly from the \\hi\\ and \\htwo\\ intensity maps. The flux is computed out to two optical radii \\rtf\\ for each of the non-interacting galaxies shown in Figure \\ref{fig3}. The solid line indicates unity and the two dashed lines a factor of two scatter. The plot thus shows that Equation \\ref{mgas} allows us to predict the total gas mass of a disk galaxy within a factor of two uncertainty.} \\vspace{0.59cm} \\label{fig4} \\end{figure*}  We examine this conclusion in Figure \\ref{fig4}, which shows the total gas mass of the galaxies in Figure \\ref{fig3} as measured in the respective \\hi\\ and \\htwo\\ intensity maps, compared to the predictions made using Equation \\ref{mgas} (where the integration is carried out to 2$\\times$\\rtf). We estimate the uncertainty on the total gas masses derived from the intensity maps (x-axis) from the calibration uncertainties of the THINGS and HERACLES data \\citep[$5\\%$ and $20\\%$, respectively,][]{walter08, leroy09}. We thus adopt $20\\%$ as the overall uncertainty. For the masses derived from the fit (y-axis), we estimate the uncertainty from error propagation (from Equation \\ref{mgas}). We take into account the uncertainties on the fit (slope and intercept), \\rtf\\ \\citep[estimated from][]{paturel91} and ${\\Sigma_{\\rm trans}}$ \\citep[from][]{leroy08}. This results in an uncertainty estimate of $\\sim40\\%$ for the mass prediction. Both uncertainties are indicated as error bars in Figure \\ref{fig4}.  Figure \\ref{fig4} shows that this equation offers a robust way to predict the total gas mass of disk galaxies to within a factor of two (indicated by the dashed lines, the solid line indicates unity). The mean ratio of predicted versus measured total gas mass is $1.17$ (with standard deviation $0.36$), which underlines the good correspondence between prediction and data. This close correspondence is rooted in the tightly constrained fit in Figure \\ref{fig3}, which falls almost directly on the means in the individual radius bins.  Based on observed metallicity gradients across galaxy disks \\citep[e.g.,][]{moustakas10}, one might speculate about a similarly radially varying CO-to-\\htwo\\ conversion factor. Recent work indeed suggests that systematic variations of \\xco\\ with radius (albeit not necessarily driven primarily by metallicity) occur in at least some of the galaxies studied in this paper: in these cases, \\xco\\ is observed to increase with increasing radius (Sandstrom et al., in prep.). A low value for XCO in galaxy centers seems to be a more common feature (Sandstrom et al., in prep.). The former effect would lead to shallower \\htwo\\ profiles , while the latter might in fact explain the observed apparent \\htwo\\ excesses in galaxy centers, possibly even giving rise to truly exponential \\htwo\\ profiles including the centers.  One of the implications of a universal gas profile is that there appears to be some mechanism that keeps the relationship constant in the face of star formation in these galaxies.  For example, \\citet[]{bigiel08,leroy08, bigiel11} show that the depletion time for the molecular gas in disk galaxies is $\\sim 2 \\times 10^9$\\,yr. The region of Figure 3 where $\\frac{\\Sigma_{\\rm gas}}{\\Sigma_{\\rm trans}}>1$ (roughly for $r< 0.45\\times\\rtf$) is the region dominated by molecular gas and that gas will therefore be exhausted in about 15\\% of a Hubble time. In order to keep the profiles self-similar, this implies that either new gas comes from outside the disk and falls preferentially in the central regions, an unlikely occurrance, or that gas flows through the disk to make up the gas lost to star formation, which takes place preferentially within $r< 0.45\\times\\rtf$. In either case, it is difficult to understand why the total gas profiles would be so self-similar.  For the infall case, one expects such infall to be sporadic, possibly occurring in the form of small galaxies merging with the bigger disk galaxies in our sample. For the inflow case, Figure \\ref{fig3} encompasses a variety of galaxy types and morphologies, and it is hard to see how inflow could be so closely regulated to produce the tight observed relationship.  As cosmological simulations of galaxy evolution including gas become ever more refined, it will become necessary for these simulations to reproduce the universal relationship when carried out to the present epoch.  Because of the limitations in making observations of the atomic gas to higher redshift with present day instrumentation, it will be difficult in the near future to extend the work presented here to normal disk galaxies at significantly higher redshifts.  It will also be a challenge to extend this work to lower metallicity systems, where CO emission becomes an increasingly poor tracer of the molecular hydrogen \\citep[e.g.,][]{bolatto11}. On the other hand, \\rtf\\ can be determined in many galaxies to higher redshifts and Equation 2 could be used to estimate M$_{\\rm gas}$. As it becomes possible to measure {\\it total} neutral gas masses to higher redshifts with ALMA, the JVLA, and Arecibo, direct comparisons of the predicted and measured total mass of neutral gas can be extended both to higher redshift and to a larger sample of local galaxies providing good tests of the universality of the neutral gas profile presented in this paper.  The universal gas profile we have described seems to be a fundamental property of normal disk galaxies at $z=0$. We have, however, only probed galaxies of near solar metallicity and none of the galaxies in our final sample are dwarfs. Also, \\citet{young11} and \\citet{serra12} have shown recently that a surprisingly large fraction of early type galaxies, i.e., ellipticals and lenticulars, contain large amounts of atomic and molecular gas.  Whether these galaxies obey the same universal gas profile is still to be determined."
    },
    "1208/1208.0302_arXiv.txt": {
        "abstract": "We present a new computational approach to the inversion of solar photospheric Stokes polarization profiles, under the Milne-Eddington model, for vector magnetography.  Our code, named {\\sc genesis} ({\\sc gene}tic {\\sc s}tokes {\\sc i}nversion {\\sc s}trategy), employs multi-threaded parallel-processing techniques to harness the computing power of graphics processing units ({\\sc gpu}s), along with algorithms designed to exploit the inherent parallelism of the Stokes inversion problem.  Using a genetic algorithm ({\\sc ga}) engineered specifically for use with a {\\sc gpu}, we produce full-disc maps of the photospheric vector magnetic field from polarized spectral line observations recorded by the Synoptic Optical Long-term Investigations of the Sun ({\\sc solis}) Vector Spectromagnetograph ({\\sc vsm}) instrument.  We show the advantages of pairing a population-parallel genetic algorithm with data-parallel {\\sc gpu}-computing techniques, and present an overview of the Stokes inversion problem, including a description of our adaptation to the {\\sc gpu}-computing paradigm.  Full-disc vector magnetograms derived by this method are shown, using {\\sc solis/vsm} data observed on 2008 March 28 at 15:45 {\\sc ut}.\\footnote{Full-resolution versions of the images in this paper are available in the journal- and electronic-version.} ",
        "introduction": "\\label{sec:intro} Since \\citet{hale:1908} first inferred the presence of magnetic fields on the sun by observing the Zeeman-induced separation of the components of magnetically-sensitive spectral lines, the reliable determination of vector fields from spectral line observations has played a pivotal role in diagnosing solar magnetism.  The Stokes vector, ${\\bf I}_{\\lambda}=[I_{\\lambda},Q_{\\lambda},U_{\\lambda},V_{\\lambda}]^{T}$, whose elements are linear combinations of measured polarization intensities at a wavelength $\\lambda$, provide a convenient way to observe and characterize the effects of magnetic fields on the absorbing medium.  \\citet{unno:1956} first described the Zeeman-induced attenuation of the Stokes vector by magnetic fields in the framework of the radiative transfer equations neglecting magneto-optical effects.  The solution was later generalized to include magneto-optical (Faraday rotation) effects by \\citet{rachkovsky:1962}.  Much information can be directly extracted from the observed Stokes vectors themselves; \\citet{rees:1979} developed a center-of-gravity technique for estimating the longitudinal component of the magnetic field from Stokes $I$ \\& $V$ observations, and \\citet{ronan:1987} described an integral method for recovering the longitudinal and transverse fields from integrated Stokes linear and circular polarization profiles.  Furthermore, several convenient weak-field calculation ``recipes'' can be found in \\citet{landideglinnocenti:1994}.  The unique specification of the magnetic and thermodynamic state of the solar photosphere solely from observations of the Stokes vector is classified as an ``inverse'' problem.  These types of problems are often computationally complex and may be ill-conditioned.  Conversely, the forward-modeling of the Stokes vector emergent from an assumed model atmosphere is trivial.  This asymmetry can be exploited to interpret the observed Stokes vector within the framework of a magnetic model atmosphere, by tuning its configuration to minimize (in a least-squares sense) the model deviation from the observed Stokes vector.  We collectively refer to all such solution procedures as ``inversion methods.'' \\citet{auer:1977} developed a now traditional Stokes inversion method based on the Levenberg-Marquardt ({\\sc l-m}) algorithm \\citep{levenberg:1944,marquardt:1963}, which was subsequently extended and improved upon by \\citet{skumanich:1987}.  This optimization approach performs a nonlinear least-squares fit of a Milne-Eddington model to the observations, returning the set of atmospheric parameters that describe the polarized spectral line.  \\citet{ruizcobo:1992} extended the optimization method by creating an inversion code called {\\sc sir} ({\\sc s}tokes {\\sc i}nversion based on {\\sc r}esponse functions), which performs the inversion of observed profiles while simultaneously inferring the stratification of the model parameters with optical depth.  More recently, artificial intelligence methods and pattern recognition approaches have been gaining ground in the computational methods of spectropolarimetric analysis.  \\citet{carroll:2001} proposed a technique based on an artificial neural network ({\\sc ann}), whereby a large database of Stokes profiles (either synthetic or pre-inverted by some other means) is used to train the {\\sc ann} to recognize the functional relationship between the model parameters and the spectral lines in the training set.  Once suitably trained, the network can generalize the relationship to other samples not explicitly included in the training set.  \\citet{rees:2000} developed a technique, based on the Principal Component Analysis ({\\sc pca}) formulated by \\citet{pearson:1901}, to decompose a line profile into their so-called {\\it eigenprofiles}.  The eigenvalues associated with these eigenprofiles define a point in the model manifold, from which the associated model parameters can be calculated by interpolation over the training set.  This method has been subsequently used on real observations by \\citet{socasnavarro:2001} and \\citet{eydenberg:2005}.  This paper introduces a new approach to the synthesis and inversion of spectral lines, based on graphics processing units ({\\sc gpu}s), which can rapidly calculate full-disc vector magnetic fields at the photospheric level.  The technique is based on the combination of a highly-parallel genetic algorithm with a computing architecture well suited to exploit the many levels of parallelism in the Stokes inversion problem.  The remainder of this paper is organized as follows.  Section \\ref{sec:data} presents the observational data used in this work.  Sections \\ref{sec:genetic} and \\ref{sec:cuda} outline our genetic algorithm optimization engine and {\\sc gpu}-programming techniques, respectively.  Details of our implementation of Stokes inversion under the assumption of a Milne-Eddington atmosphere are given in Section \\ref{sec:implementation}.  Results from the analysis presented in Section \\ref{sec:results}.  Finally, we offer some outlooks on the future implementation of our method in the vector field data reduction pipeline for the Vector Spectromagnetograph ({\\sc vsm}), the scanning spectropolarimeter instrument package in operation as part of the Synoptic Optical Long-term Investigations of the Sun ({\\sc solis}) telescope, located at the National Solar Observatory atop Kitt Peak, AZ. ",
        "conclusions": "\\label{sec:conclusions} We have described a novel computational approach to the inference of photospheric vector magnetic fields from observations of the Stokes polarization profiles.  Our new inversion code, named {\\sc genesis}, is capable of quickly producing full-disc spatial maps of the magnetic structure of the solar photosphere observed by the {\\sc solis/vsm} instrument located atop Kitt Peak at the National Solar Observatory, in a fraction of the time required by a similar serial technique.  The inversion code is capable of recovering magnetic fields with uncertainties on the order of 0.5\\%, with errors in the field orientation of a few degrees.  Fill-fractions are recovered with uncertainties on the order of 2\\%.  Currently, the code is only capable of inverting a single line at a time, though we plan to investigate the extension to the simultaneous inversion of both lines of the Fe {\\sc i} $\\lambda$6302 multiplet.  We have shown the technique to be amenable to the reduction and analysis of large volumes of spectropolarimetric data.  To this end, we are currently investigating the the assimilation of the {\\sc gpu} hardware and specialized {\\sc cuda}-based algorithms into the {\\sc solis/vsm} vector field pipeline. Our long-term goals for this work are to provide near real-time vector magnetic fields to the scientific community.  Increasing the cadence of {\\sc solis/vsm} vector data products will also allow us to support and complement observations taken by the Solar Dynamics Observatory ({\\sc sdo}) Helioseismic and Magnetic Imager ({\\sc hmi}), which produces full-disc vector magnetograms at $4096 \\times 4096$ resolution with a cadence of approximately 12 minutes.  The {\\sc gpu} programming paradigm is highly scalable; a compiled {\\sc cuda} application can execute on any {\\sc cuda}-capable device, subject to hardware limitations.  The thread scheduler will automatically allocate the appropriate number of thread blocks to the stream processors.  Coupled with the Message Passing Interface ({\\sc mpi}) to parallelize over scanlines (i.e. independent scanlines are inverted independently, utilizing their own distinct {\\sc gpu}), this could lead to incredibly fast (i.e. near-realtime or realtime) full-disc inversions with a modest number of {\\sc cpu-gpu} pairs. Improvements in {\\sc gpu} hardware are steadily advancing; the Fermi architecture is the recent successor to the Tesla architecture, offering up to 512 {\\sc cuda} cores, larger-capacity memory banks, and increased floating-point performance.  With the next-generation Kepler architecture on the horizon, the prospects for accelerated solar data processing are indeed promising.  Finally, as {\\sc gpu} computing matures, we expect to extend this approach to better hardware, with a cautious eye toward spectropolarimetric analyses of data recorded by the upcoming Advanced Technology Solar Telescope ({\\sc atst}).  The volume of data expected from this next-generation observatory will greatly exceed that of the current generation, requiring new and faster techniques to properly handle and reduce the observations in a timely manner.  We feel the integration of {\\sc gpu}-accelerated data-reduction techniques will be key for the analysis of such large datasets, and may make available important (near) real-time information on the photospheric vector magnetic field to the space-weather forecasting community."
    },
    "1208/1208.5461_arXiv.txt": {
        "abstract": "\\noindent We present deep optical and near--infrared (NIR) $UBVRIHKs$ imaging data for $24$ blue compact galaxies (BCGs). The individual exposure times are on average $\\sim40$ minutes in the optical ($B$) and $\\sim90$ minutes in the NIR, but on occasion up to $\\sim5$ hours for a single target and filter, observed on $2.5,~3.5,~8.2$--m telescopes. The sample contains luminous dwarf and intermediate--mass BCGs which are predominantly metal--poor, although a few have near--solar metallicities. We have analyzed isophotal and elliptical integration surface brightness and color profiles, extremely deep ($\\mu_B\\lesssim29$ mag arcsec${}^{-2}$) contour maps and RGB images for each galaxy in the sample, and provide a morphological classification where such is missing. We have measured the total galaxy colors, the colors of the underlying host galaxy, and the colors of the burst, and compare these to the predictions of new state--of--the--art spectral evolutionary models (SEMs) both with and without contribution by nebular emission. Separating the burst from the underlying host we find that regardless of the total luminosity the host galaxy has the properties of a low surface brightness (LSB) dwarf with $M_B\\gtrsim-18$. For a number of galaxies we discover a distinct LSB component dominant around and beyond the Holmberg radius. For the specific case of ESO400--43A\\&B we detect an optical bridge between the two companion galaxies at the $\\mu_V\\sim28$th mag arcsec${}^{-2}$ isophotal level. Synthetic disk tests are performed to verify that we can trace such faint components with negligible errors down to $\\mu_B=28$ mag arcsec${}^{-2}$ and $\\mu_K=23$ mag arcsec${}^{-2}$. By examining the structural parameters (central surface brightness $\\mu_0$ and scale length $h_r$) derived from two radial ranges typically assumed to be dominated by the underlying host galaxy, we demonstrate the importance of sampling the host well away from the effects of the burst. We find that $\\mu_0$ and $h_r$ of the BCGs host deviate from those of dwarf ellipticals (dE) and dwarf irregulars (dI) solely due to a strong burst contribution to the surface brightness profile almost down to the Holmberg radius. Structural parameters obtained from a fainter region, $\\mu_B=26$--$28$ mag arcsec${}^{-2}$, are consistent with those of true LSB galaxies for the starbursting BCGs in our sample, and with dEs and dIs for the BCGs with less vigorous star formation. ",
        "introduction": "\\noindent Blue compact galaxies (BCGs) are gas--rich star--forming low redshift galaxies with low metallicities~\\citep[for a review see][]{2000A&ARv..10....1K}. Their star formation rates (SFRs) can be as high as $\\sim20~M_\\odot~yr^{-1}$~\\citep[e.g. Haro 11,][]{2007MNRAS.382.1465H} and their \\HI gas masses range from $M(H\\texttt{I})\\sim10^7$ to a few tens of $10^8 M_\\odot$~\\citep[e.g.][]{1995ApJS...99..427T,1999A&AS..139....1T,2000A&A...361...19S}. BCGs are said to be starbursting when they will exhaust their gas supply in less than a Hubble time. The metallicities measured in these galaxies are low~\\citep[e.g.][]{1994ApJ...420..576M,1994MNRAS.270...35M,1999ApJ...511..639I}, to extremely low~\\citep{2005ApJ...632..210I}, once leading researchers to propose that BCGs are truly young galaxies~\\citep{1972ApJ...173...25S}, experiencing their first episode of star formation in a primordial gas environment. There is mounting evidence to the contrary, however, and ``first--burst'' claims have been challenged using broadband colors and morphology~\\citep[e.g.][]{1996A&A...314...59P,1997MNRAS.286..183T,2001ApJS..133..321C,2001ApJS..136..393C,2002A&A...390..891B}, absorption features in the spectra~\\citep[e.g.][]{2004A&A...423..133W}, through the direct detection of globular clusters associated with the BCG~\\citep{1998A&A...335...85O}, and for near--by targets through the detection of luminous red giants and asymptotic giant branch (AGB) stars in color--magnitude diagrams~\\citep[e.g.][]{1999AJ....118..302A,2000ApJ...535L..99O}. In most cases the relative strength of the starburst can be so high that it completely dominates the light output of the galaxy, an obstacle which has been countered by deeper optical imaging data and observations in the near infra-red (NIR) regime~\\citep[e.g.][]{2003A&A...410..481N,2003ApJ...593..312C} and an older population has been revealed. In other cases, e.g. SBS0335--052 E\\&W, the burst can be spatially dominating over the galaxy even for extremely deep imaging, making it difficult to directly detect a possible old population. Hence, the youth hypothesis persists for these targets~\\citep[e.g.][]{1997ApJ...489..623T,2005ApJ...632..210I} but has been challenged~\\citep{2001A&A...371..429O}.\\\\  \\noindent Nevertheless, the emerging consensus seems to be that BCGs are galaxies with an underlying old stellar population, often referred to as 'host', on which the current starburst is superposed. The host population is distinctly different in terms of optical and NIR colors~\\cite[e.g.][]{2002A&A...390..891B,2003A&A...410..481N,2003ApJ...593..312C,2005mmgf.conf..355B,2005ApJS..157..218C,2009A&A...501...75A}, and appears to be smoothly distributed in a physical structure reminiscent of a disk with regular isophotes~\\citep{1983MitAG..58..108L,1983ApJ...268..667T,1996A&A...314...59P,1999A&A...341..725S}. The discovery of this host population invariably led to a relaxation in the definition of this group of galaxies. The ``dwarf'' qualifier in the original naming convention of ``blue compact dwarf'' is often dropped by some authors to the benefit of the more inclusive ``galaxy'' term. There is a large fraction of BCGs which are too massive and bright to be dwarfs, with $M_B$ as high as $-20$, but they too exhibit starbursts in a gas--rich, metal--poor environment, and a similar morphological structure as the dwarfs -- an extended, often regular, disk--like structure hosts knots of star formation (SF), either centrally located or dispersed in seemingly random fashion around the galaxy. The name is still misleading since, as pointed out in~\\citet{2003ApJS..147...29G}, some blue compact galaxies are neither blue, nor compact. Including the host, the requirements of~\\citet{1981ApJ...247..823T} that a BCG must have an apparent optical size $r_{25}\\lesssim1$ kpc in diameter can now only be met for some very compact sources, like HL293B. Further, the ``blue'' condition of $B-V\\lesssim0.3^{\\textrm{m}}$~\\citep{1986A&AS...64..469B} can now only be applied to the starbursting region since inclusion of the host shifts the total $B-V$ color redwards of this (somewhat arbitrary) limit.\\\\  \\noindent Despite such diversity all BCGs have at least one thing in common. They are a subgroup of emission line galaxies, well separated from Seyferts and quasars, with spectra reminiscent of \\HII regions, only with a strong stellar continuum from the composite populations of the galaxy~\\citep[e.g.][]{1983ApJ...273...81K,1994ApJ...420..576M,1989ApJS...70..447S}. Because of their \\HII region--like spectrum they are often referred to as \\HII galaxies. The SFR can drastically vary from BCG to BCG, with some extreme outliers like \\emph{Haro 11} having a $SFR\\gtrsim18~M_\\odot~yr^{-1}$, which it can only sustain for $\\lesssim50$ Myr~\\citep{2000A&A...359...41B}. While not equally vigorous, the SF activity in most BCGs must similarly proceed in short--lived bursts, lasting from a few tens of Myr to $1$ Gyr. The duration of the typical burst has been constrained both from observed \\HI masses, which are generally too low to sustain the current SFRs for a Hubble time, and from measured metallicities, which are also too low for the interstellar medium to have undergone enrichment by constant star formation. The latter argument can be partially circumvented if one allows for the outflow and loss of metals to the intergalactic medium due to supernovae winds, or for the inflow of pristine gas, which would dilute the metallicity. The argument of too low metallicity cannot be completely discarded, however, as there are indications that gas--rich galaxies are more likely to be closed systems~\\citep{2006ApJ...636..214V}.\\\\  \\noindent Some BCGs are clearly in various stages of merging, or undergoing tidal interaction. They can have very irregular morphologies, with complicated, almost chaotic kinematics, making it difficult to quantify the masses and ages of the current burst. Other BCGs show an otherwise regular elliptical disk with a central nuclear starburst, or sporadic star--forming knots offset from the center, a.k.a. \\HII hotspots~\\citep{1989ApJS...70..447S}, and it is not clear whether galaxies with such morphologies also are subject to the same SF trigger. \\HI observations have shown that while the majority of BCGs are associated with massive \\HI clouds, most likely fueling the ongoing star formation, they are generally isolated in the sense that they lack massive companion galaxies~\\citep{1991A&A...241..358C,1993AJ....106.1784C,1995ApJS...99..427T,1999A&AS..139....1T}. However, \\HI companions without a confirmed optical counterpart do appear to be common~\\citep{1997ApJ...480..524T}, indicating that perhaps the companion galaxy is a faint gas--rich dwarf. This makes mergers the leading candidates for SF triggers since tidal interactions are unlikely to be the primary SF triggers in dwarf galaxies~\\citep{2004MNRAS.349..357B}. For a starburst to take place the merger would not have to be major but it would have to be wet, meaning at least one of the merging galaxies would have to be gas--rich.\\\\  \\noindent Ideally, one would like to study the SF as it occurs for the first time in the protogalaxies in the distant Universe. It is thought that observing such pristine SF activity would allow us to determine what the necessary trigger conditions for the onset of star formation are. Presently, we are unable to directly study the stellar populations of high--redshift galaxies in any great detail since we are constrained by cosmological effects and small intrinsic sizes of the targets, and the comparatively poor resolution of our currently available instrumentation. Because of their low metallicities, high SFRs, and their lack of spiral density waves, BCGs are thought to be reminiscent of the young galaxies in the universe, in which the star formation must have proceeded under similar metal--poor gas--rich conditions. Being much closer, BCGs are available for detailed photometric and spectroscopic studies.\\\\  \\noindent Considering the degree of complication inherent to the analysis of such a heterogeneous group of galaxies, one needs to examine the properties of the BCGs from several angles: the behavior of the spectrum in the optical and NIR regimes on a large scale (broadband imaging), the individual line emissions (typically $H_\\alpha$ and $[OIII]$), and the kinematics of the gas and the stars. We have undertaken a study of BCGs that combines optical and NIR broadband and narrowband imaging~\\citep[this work,][]{Paper2,Paper3}, Fabry--Perot interferometry (\\\"Ostlin et al. 2012 in prep., Marquart et al. 2012 in prep), and grism spectroscopy of two distinct types of star--forming galaxies - high and low luminosity BCGs. In this paper, which is the first in a series, we present broadband $UBVRIHKs$ data for a sample of $24$ luminous galaxies classified as BCGs. This sample was hand--picked to contain interesting and representative cases, and is biased towards relatively luminous (median $M_B\\sim-18$ mag) galaxies. Some of the targets come from the~\\citet{2001A&A...374..800O} sample, others from the~\\citet{1989ApJS...70..447S} UM survey of \\HII galaxies. The latter were selected based on their emission line spectra to be actively star--forming if not always strongly starbursting, and to have metallicities and \\HI masses available in the literature. Additionally, a number of famous and strongly starbursting BCGs (e.g. IIZw40, MK930, Haro11) were included in the project because they lacked equally deep observations and yet presented interesting targets to have in our analysis. Haro 11 was already presented in~\\citet{2010MNRAS.405.1203M}, so it is not presented here but will be included in our future analysis. In a consecutive paper~\\citep{Paper2} we present $UBVRIHKs$ broadband imaging data for a volume--limited sample of $21$ low luminosity emission line galaxies. These constitute a well defined complete sample of star forming galaxies, representative of the galaxy population in the local Universe. In total we have $46$ BCG and BCG--like targets in $7$ broadband and $2$ narrowband filters. In the optical, imaging observations of depth comparable to ours can be found in the literature for a number of galaxies~\\citep{1996A&AS..120..207P,1999A&AS..138..213D,2001ApJS..133..321C,2002A&A...390..891B}. In the NIR, however, there is no sample of BCGs with as deep individual exposure times as ours~\\citep[compare e.g.][]{2003A&A...410..481N,2003ApJ...593..312C}. In~\\citet{Paper3} we combine both samples to juxtapose the high and low luminosity BCGs and analyse any differences in their host populations. In the latter paper we also present the narrowband data from both samples as well as derived estimates of the ages, masses and star formation rates through SED fitting. We have assumed $H_0=73~km^{-1}s^{-1}Mpc^{-1}$ throughout this paper. \\\\  \\noindent The layout of this paper is as follows: \\S~\\ref{data} introduces the observations, reduction and calibration of the data. \\S~\\ref{sbsection}, \\S~\\ref{contourplots}, and \\S~\\ref{rgbimages} deal with the generation of surface brightness profiles, contour plots and three--color RGB images, respectively. \\S~\\ref{integrsurfphot} and \\S~\\ref{colors} explain our approach to obtaining total luminosities and total colors for the galaxies. \\S~\\ref{strucparam} illustrates how we obtain structural parameters for population separation. Characteristics of individual galaxies and comparison of integrated colors to stellar evolutionary models can be found in \\S~\\ref{individ}. Finally, discussion and concluding remarks are in \\S~\\ref{discuss} and \\S~\\ref{conclude}.  \\section[]{Observations}\\protect\\label{data}  \\noindent The data consist of optical and near infra-red (NIR) broadband imaging in the $UBVRIHK$ filters. They were obtained during the period $2001$--$2007$ with ALFOSC (Nordic Optical Telescope, NOT), EMMI (European Southern Observatory New Technology Telescope, ESO NTT), and FORS (ESO Very Large Telescope, VLT) in the optical, and with NOTCAM (NOT) and SOFI (ESO NTT) in the NIR. A log of observations is shown in Table~\\ref{nedtable} and the individual exposure times for each filter in Table~\\ref{exptable}. Since both the sample presented here and the volume limited sample of~\\citet{Paper2} were reduced with the same pipelines, the description of the reductions and calibration that follows here applies to both. We will give the details in this paper and will not repeat them in~\\citet{Paper2}.  \\setcounter{table}{0} \\begin{table} \\begin{minipage}{70mm} \\tiny \\caption{Total integration times for the sample. All times are given in minutes and converted to the framework of a 2.56 meter telescope where needed. The values are for observations in a single filter, e.g. only $SOFI~Ks$, and not $SOFI~Ks+NOTCAM~Ks$.}\\protect\\label{exptable} \\begin{tabular}{|lccccccc|} % \\hline &U&B&V&R&I&H&Ks\\\\\\hline\\hline ESO185--13&30&38&9&7&38&67&137\\\\\\hline ESO249--31&&38&19&19&19&58&117\\\\\\hline ESO338--04&&9&19&29&125&61&352\\\\\\hline ESO400--43&&&77&&148&61&314\\\\\\hline ESO421--02&30&38&9&19&19&19&114\\\\\\hline ESO462--20&30&38&9&&&73&139\\\\\\hline HE2--10&51&38&19&&&9&121\\\\\\hline HL293B&26&45&25&20&40&25&64\\\\\\hline IIZW40&10&40&20&&40&12&$96^\\dagger$\\\\\\hline MK600&20&40&23&20&40&24&119\\\\\\hline MK900&20&40&25&20&40&21&32\\\\\\hline MK930&20&40&25&30&40&37&127\\\\\\hline MK996&20&40&20&20&26&15&75\\\\\\hline SBS0335--052E\\&W&&103&&&45&94&$298^\\ddagger$\\\\\\hline TOL0341--407&30&29&9&&38&67&117\\\\\\hline TOL1457--262\\emph{I}\\&\\emph{II}&&66&76&&&&125\\\\\\hline UM133&20&53&23&20&40&22&53\\\\\\hline UM160&20&40&26&20&40&96&56\\\\\\hline UM238&20&40&23&20&40&21&96\\\\\\hline UM417&20&40&23&20&40&18&58\\\\\\hline UM448&&40&40&9&58&&32\\\\\\hline UM619&40&40&40&9&58&32&$186^\\dagger$\\\\\\hline \\end{tabular} \\medskip ~\\\\ $\\dagger$ -- SOFI, $\\ddagger$ -- NOTCAM \\end{minipage} \\end{table} \\begin{center} \\begin{table*} \\begin{minipage}{150mm} \\tiny \\caption{Log of the observations. Heliocentric redshift and cosmology--corrected luminosity distances from \\emph{NED}.}\\protect\\label{nedtable} \\begin{tabular}{@{}|lllllll|@{}} \\hline Galaxy&Ra~Dec~(J2000)&Redshift&$\\textrm{D}~[\\textrm{Mpc}]$&Year&Instrument&Filters\\\\\\hline ESO185--13&19h45m00.5s&0.018633&76.3&2001&EMMI&Bb\\#605,V\\#606,I\\#610\\\\ &-54d15m03s&&&&SOFI&Ks\\#13\\\\ &&&&2002&FORS1&U\\_BESS,R\\_BESS\\\\ &&&&&SOFI&H\\#12\\\\\\hline ESO249--31&03h55m46.1s &0.002829&10.8&2001&EMMI&Bb\\#605,V\\#606,R\\#608,I\\#610\\\\ &-42d22m05s&&&&SOFI &Ks\\#13\\\\ &&&&2002 &SOFI&H\\#12\\\\\\hline ESO338--I04&19h27m58.2s &0.009453&37.4&2001&EMMI&R\\#608\\\\ &-41d34m32s&&&&EMMI&Bb\\#605,V\\#606\\\\ &&&&2005&EMMI&I\\#610\\\\ &&&&2005&SOFI&Ks\\#13\\\\ &&&&2006&SOFI&Ks\\#13,H\\#12\\\\\\hline ESO400--G043&20h37m41.9s&0.019680&79.2&2005&EMMI&V\\#606,I\\#610\\\\ &-35d29m08s&&&2006&SOFI&KS\\#13,H\\#12\\\\\\hline ESO421--02&04h29m40.1s&0.003149&12.4&2001&EMMI&Bb\\#605,V\\#606,R\\#608,I\\#610\\\\ &-27d24m31s&&&&SOFI&Ks\\#13\\\\ &&&&2002&FORS1&U\\_BESS\\\\ &&&&&SOFI&H\\#12\\\\\\hline ESO462--IG020&20h26m56.8s&0.019782&79.5&2001&EMMI&Bb\\#605,V\\#606\\\\ &-29d07m01s&&&&SOFI&Ks\\#13\\\\ &&&&2002&FORS1&U\\_BESS\\\\ &&&&&SOFI&H\\#12\\\\\\hline HE2--10&08h36m15.1s&0.002912&15.8&2001&EMMI&Bb\\#605,V\\#606\\\\ &-26d24m34s&&&2002&FORS1&U\\_BESS\\\\ &&&&&SOFI&H\\#12,Ks\\#13\\\\\\hline HL293B&22h30m36.798s&0.005170&16.3&2001&ALFOSC--FASU&B\\#74,V\\#75,R\\#76,I\\#12\\\\ &-00d06m36.99s&&&2002&ALFOSC--FASU&U\\#7\\\\ &&&&2003&NOTCAM&Ks\\#207,H\\#204\\\\\\hline IIZW40&05h55m42.6s&0.002632&11.8&2001&ALFOSC--FASU&B\\#74,I\\#12\\\\ &+03d23m32s&&&2002&ALFOSC--FASU&U\\#7\\\\ &&&&&NOTCAM&Ks\\#207\\\\ &&&&&SOFI&Ks\\#13\\\\ &&&&2003&ALFOSC--FASU&V\\#75\\\\ &&&&&NOTCAM&H\\#204\\\\\\hline MK600&02h51m04.6s&0.003362&10.9&2001&ALFOSC--FASU&B\\#74,I\\#12\\\\ &+04d27m14s&&&2002&ALFOSC--FASU&U\\#7,V\\#75,R\\#76\\\\ &&&&&NOTCAM&Ks\\#207\\\\ &&&&2003&NOTCAM&H\\#204\\\\\\hline MK900&21h29m59.6s&0.003843&11.2&2001&ALFOSC--FASU&B\\#74,V\\#75,R\\#76,I\\#12\\\\ &+02d24m51s&&&2002&ALFOSC--FASU&U\\#7\\\\ &&&&2003&NOTCAM&Ks\\#207,H\\#204\\\\\\hline MK930&23h31m58.3s&0.018296&71.4&2001&ALFOSC--FASU&B\\#74,V\\#75,I\\#12\\\\ &+28d56m50s&&&2002&ALFOSC--FASU&U\\#7,R\\#76\\\\ &&&&2003&NOTCAM&Ks\\#207,H\\#204\\\\\\hline MK996&01h27m35.5s&0.005410&18.2&2001&ALFOSC--FASU&I\\#12\\\\ &-06d19m36s&&&2002&ALFOSC--FASU&U\\#7\\\\ &&&&&SOFI&Ks\\#13\\\\ &&&&2003&ALFOSC--FASU&B\\#74,V\\#75,R\\#76\\\\ &&&&&NOTCAM&H\\#204\\\\\\hline SBS0335--052E\\&W&03h37m44.0s&0.013486&54&2001&ALFOSC--FASU&B\\#74,I\\#12\\\\ &-05d02m40s&&&2002&SOFI&Ks\\#13\\\\ &&&&2003&NOTCAM&KS\\#207,H\\#204\\\\\\hline TOL0341&03h42m49.4s&0.014827&60.6&2001&EMMI&Bb\\#605,V\\#606,I\\#610\\\\ &-40d35m56s&&&2002&FORS1&U\\_BESS\\\\ &&&&&SOFI&H\\#12,Ks\\#13\\\\\\hline TOL1457--262\\emph{I}\\&\\emph{II}&15h00m27.9s&0.016812&72.8&2004&ALFOSC--FASU&B\\#74,V\\#75\\\\ &-26d27m02s&&&2005&SOFI&Ks\\#13\\\\\\hline UM133&01h44m41.3s&0.005414&18.3&2001&ALFOSC--FASU&B\\#74,I\\#12\\\\ &+04d53m26s&&&2002&ALFOSC--FASU&U\\#7,V\\#75,R\\#76\\\\ &&&&2003&NOTCAM&Ks\\#207,H\\#204\\\\\\hline UM160&23h24m19.8s&0.008006&28&2001&ALFOSC--FASU&B\\#74,V\\#75,I\\#12\\\\ &-00d07m01s&&&2002&ALFOSC--FASU&U\\#7,R\\#76\\\\ &&&&&SOFI&H\\#12,Ks\\#13\\\\\\hline UM238&00h24m42.3s&0.014230&54.2&2001&ALFOSC--FASU&B\\#74,I\\#12\\\\ &+01d44m02s&&&2002&ALFOSC--FASU&U\\#7,V\\#75,R\\#76\\\\ &&&&&SOFI&Ks\\#13\\\\ &&&&2003&NOTCAM&H\\#204\\\\\\hline UM417&02h19m30.2s&0.009000&33.8&2001&ALFOSC--FASU&B\\#74,I\\#12\\\\ &-00d59m11s&&&2002&ALFOSC--FASU&U\\#7,V\\#75,R\\#76\\\\ &&&&&NOTCAM&Ks\\#207\\\\ &&&&2003&NOTCAM&H\\#204\\\\\\hline UM448&11h42m12.4s&0.018559&74.6&2004&ALFOSC--FASU&B\\#74,V\\#75\\\\ &+00d20m03s&&&2005&ALFOSC--FASU&U\\#7\\\\ &&&&&EMMI&R\\#608,I\\#610\\\\ &&&&2007&NOTCAM&Ks\\#207\\\\\\hline UM619&13h52m44.7s&0.015494&63.1&2004&ALFOSC--FASU&B\\#74,V\\#75\\\\ &+00d07m53s&&&2005&ALFOSC--FASU&U\\#7\\\\ &&&&&EMMI&R\\#608,I\\#610\\\\ &&&&&NOTCAM&H\\#204\\\\ &&&&&SOFI&Ks\\#13\\\\ &&&&2006&SOFI&Ks\\#13\\\\ &&&&2007&NOTCAM&Ks\\#207\\\\\\hline \\hline \\end{tabular} \\end{minipage} \\end{table*} \\end{center} \\begin{center} \\begin{figure*} \\begin{minipage}{150mm} \\centering \\includegraphics[width=15cm,height=18cm]{fig1.ps} \\caption{\\textbf{ESO185-13}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $22.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\protect\\label{datafig} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig2.ps} \\contcaption{\\textbf{ESO249-31}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $22.0$, $24.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$, $25$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,V,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig3.ps} \\contcaption{\\textbf{ESO338-04}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $22.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,V,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig4.ps} \\contcaption{\\textbf{ESO400-43}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $V$ band. Isophotes fainter than $23.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $V,I,H$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig5.ps} \\contcaption{\\textbf{ESO421-02}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig6.ps} \\contcaption{\\textbf{ESO462-20}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $22.5$ are smoothed with a boxcar median filter of size $5$ pixels. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,V$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig7.ps} \\contcaption{\\textbf{HE2-10}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $22.5$, $24.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$, $25$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,V$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig8.ps} \\contcaption{\\textbf{HL293B}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig9.ps} \\contcaption{\\textbf{IIZw40}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $22.5$, $25.0$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,V,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage   \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig10.ps} \\contcaption{\\textbf{MK600}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.0$, $25.0$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig11.ps} \\contcaption{\\textbf{MK900}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $22.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig12.ps} \\contcaption{\\textbf{MK930}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $24.0$, $26.0$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig13.ps} \\contcaption{\\textbf{MK996}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.0$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage   \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig14.ps} \\contcaption{\\textbf{SBS0335-052 EAST}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $26.0$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,I,K$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*} \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig15.ps} \\contcaption{\\textbf{SBS0335-052 WEST}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $26.0$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,I,K$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig16.ps} \\contcaption{\\textbf{TOL0341-407}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig17.ps} \\contcaption{\\textbf{TOL1457--262 \\emph{I}}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $24.0$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,V,K$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig18.ps} \\contcaption{\\textbf{TOL1457--262 \\emph{II}}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $24.0$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,V,K$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage   \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig19.ps} \\contcaption{\\textbf{UM133}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.0$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig20.ps} \\contcaption{\\textbf{UM160}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$ are smoothed with a boxcar median filter of size $5$ pixels. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig21.ps} \\contcaption{\\textbf{UM238}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage  \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig22.ps} \\contcaption{\\textbf{UM417}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $26.0$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig23.ps} \\contcaption{\\textbf{UM448}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.9$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $B,V,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*}  \\clearpage \\begin{figure*} \\begin{minipage}{150mm} \\includegraphics[width=15cm,height=18cm]{fig24.ps} \\contcaption{\\textbf{UM619}. \\textit{Left panel}: Surface brightness and color radial profiles for elliptical (open circles) and isophotal (red circles) integration. \\textit{Upper right panel}: contour plot based on the $B$ band. Isophotes fainter than $23.5$, $25.5$ are iteratively smoothed with a boxcar median filter of size $5$, $15$ pixels respectively. \\textit{Lower right panel}: A true color RGB composite image using the $U,B,I$ filters. Each channel has been corrected for Galactic extinction following \\citet{1998ApJ...500..525S} and converted to the AB photometric system. The RGB composite was created by implementing the \\citet{2004PASP..116..133L} algorithm.} \\end{minipage} \\end{figure*} \\end{center}  \\clearpage   \\subsection{Data reductions} \\noindent In order to ensure the data consistently undergo the same processing, we have reduced all raw frames, except BVR data for \\textit{ESO338--04}, with two pipelines specifically designed for this project - one for the optical and one for the NIR. This consistency ensures us that any difference in the results can not be attributed to the reduction process itself and instead reflects a true variation in properties between different galaxies. Both pipelines were already described in detail in~\\citet{2010MNRAS.405.1203M}, so here we will only present additional points of interest relevant to the reduction. \\\\  \\noindent The reductions were carried out on subsamples grouped by night and instrument. For every night and all instruments an individual sky background was fit to each reduced (i.e. bias--subtracted and flatfielded) frame \\textit{before} combining them to a final stack, and at no point in the reduction process was a ``master sky'' ever used. After stacking, a final sky fitting and subtraction was performed. Calibration in both wavelength regimes is automatic, using Landolt standard stars in the optical and Two Micron All Sky Survey (2MASS) stars in the NIR.\\\\  \\noindent \\textbf{ALFOSC}~~~Some nights on this instrument were lacking usable standard stars. However, for all such observations short $3$ minute exposures were taken during photometric nights, or the same target was observed on several occasions. It was therefore always possible to calibrate problematic nights against photometric observations of the same object.\\\\  \\noindent For $2004$ data there was a severe charge bleed in the frames of UM439, which the observers attempted to fix by varying the rotator field angle between consecutive frames of that target. This varying orientation of the images proved to be too much for the IRAF XYXYMATCH procedure, which the pipeline uses on SExtractor sources in order to align the frames. We chose to rotate all frames to the same orientation before reducing them, so there is one extra interpolation involved in the processing of these frames, as the angles were such that a simple flipping of the images in $x$, $y$ or both was not sufficient.\\\\  \\noindent \\emph{I} band frames taken with ALFOSC have pronounced fringes -- interference patterns due to the \\textit{OH} emissions coming from the upper ionosphere. If the fringes are stable throughout the night then it would be sufficient to obtain a fringe map by stacking non-aligned dithered \\emph{I} band frames and then subtract this from each individual image during reductions. This technique did not always remove the fringe pattern to our satisfaction, and residual ``hills'' and ``valleys'' could be clearly seen in the de-fringed images upon examination. Therefore we executed the pipeline in ``optical pair-subtraction mode'', using either dither or chopping mode pair-subtraction, as the data allowed. Since the fringes usually remain stable on such short time scales ($\\sim10$ minutes from the start of one exposure to the next), this always produces clearly superior results compared to the average fringe map technique, and no fringe residuals can be seen on the de-fringed image even after smoothing it with a boxcar average (the smoothing is done solely for the purpose of inspection). \\\\  \\noindent \\textbf{EMMI}~~~There was a noticeable amplifier glow at varying positions along the edges of both CCDs in most \\emph{I} band frames of the $2005$ data, which also seemed to change in time in position and intensity. If left uncorrected, this produces a significant non-flatness of the background in the stacked frame, which is impossible to remove with sky subtraction of a fitted plane. This issue is further enhanced by the fact that all four amplifiers were used during readout, each amplifier giving rise to its own varying amplifier glow. In such cases the pipeline was executed in ``optical pair-subtraction mode'', using either dither or chopping pair-subtraction, which successfully removed all amplifier glow.\\\\  \\noindent \\textbf{NOTCAM}~~~Approximately $8000$ of the raw NOTCAM images (years $2002$, $2004$ and $2005$) were lacking a world coordinate system (WCS) in the header. As the NIR pipeline relies on the WCS not only to align but also to identify 2MASS stars for calibration, we had to find a way to introduce a WCS for every individual raw frame. Obtaining a full plate solution in any kind of semi-interactive fashion for this many images is not practical. Instead we used the publicly available \\textit{astrometry.net} software~\\citep{2010AJ....139.1782L} to add the appropriate WCS to the headers. For most images the solver executed successfully using the standard indices built on the USNO-B catalog, but for a number of stubborn frames we successfully used 2MASS indices instead. Since \\textit{astrometry.net} allows the user to search for solutions using already solved nearby fields, we were able to solve any remaining frames using local indices (e.g. an entire dither sequence of images of the same object can usually be solved as long as just one of those frames solves from USNO-B or 2MASS indices).\\\\  \\noindent The Wide--Field camera of NOTCAM suffers from an optical distortion which becomes severe towards the edges of the array, and can shift the projected location of sources by several arcseconds. We have applied a correction for this using the available distortion map on the NOTCAM website. This map is from the year $2005$ and it worked well for our $2002$, $2004$, $2005$ and $2006$ NOTCAM data. We were unable to successfully remove all distortion for $2007$ data with the same distortion map. For that data we chose to instead align the images using only the central $160\\times160$ arcsec ($\\sim700\\times700$ pixels) around the galaxy. Since all targets obtained in $2007$ have a fairly small apparent size, the lingering distortion along the edges is of no consequence for our analysis.\\\\  \\noindent \\textbf{SOFI}~~~There were no significant problems during the reductions of these data. All necessary calibration frames were of good quality, and a WCS was present in all frames.\\\\  \\noindent \\textbf{FORS}~~~The FORS exposures were taken for the purpose of re-calibrating existing observations, or to fill in for a number of filters that were missing in our table of observations. There were no significant problems during the reductions.\\\\ \\begin{figure} \\includegraphics[width=8cm,height=16cm]{fig25.ps} \\caption{A typical radial error composition (here for ESO185--13). The standard deviation of the mean inside a ring ($\\sigma_{\\textrm{sdom}}$, open circles) dominates the error budget in the central regions of the galaxy. As the signal to noise ratio decreases with increasing radius the uncertainty in the sky ($\\sigma_{\\textrm{sky}}$, filled squares) becomes the dominant error source. The composite error (solid line) is included in the errorbars of all surface brightness and color profiles.}\\protect\\label{errorplot} \\end{figure} \\subsection{Photometric calibration}\\protect\\label{photometry} \\noindent All data were calibrated in the Vega photometric system. \\citet[][Northern]{1998AJ....115.2594H} and~\\citet[][Southern]{1998AJ....116.2475P} photometric standards were obtained for the NIR regime. However, we decided early on that the pipelines should calibrate the NIR data using secondary 2MASS standards on each raw frame. This way the photometric accuracy of the stacked NIR frames is nearly independent of the actual photometric conditions during the observations. For some targets the field of view contained too few visible stars, and calibration of each individual frame before stacking was then impossible. In these cases a single zero point per galaxy per night was obtained and applied to the reduced stacked images since those always have more sources than the single exposure frames.\\\\  \\noindent In the optical the data were calibrated against images of Landolt standard stars obtained during each observing night. Occasionally we ran into problems when calibrating. Some nights were clear but non-photometric, others allowed for observations only in the beginning and the end of the night, resulting in an unacceptably long time interval between galaxy and standard star observations, and yet others had too few or no standard stars available because of saturation or bad pixels. Data from nights with bad atmospheric conditions were re-calibrated using short $1$ or $3$ minute exposures, taken during photometric nights with FORS1 (VLT) or ALFOSC (NOT) in all cases where the zero point reliability was questionable.\\\\  \\noindent As we aim to reach very faint surface brightness levels, we felt it necessary to perform a sanity check on our calibration. We therefore compared the photometry of the stars in our frames to that of values taken from the Pickles stellar library in various color-color plots, both in the optical and NIR. This comparison showed no significant offset between the two,  and we are therefore confident that our photometry is reliable\\protect\\footnote{\\footnotesize These plots are available on demand. Email GM.}.\\\\  \\noindent We estimated the zero point uncertainty for the optical data using the scatter of measured zero point values for each night. For the NIR we instead measured the mean offset in the magnitudes of stars on our frames when compared to 2MASS photometry. This uncertainty is represented by $\\sigma_{zp}$ (\\S~\\ref{errorest}).\\\\ \\begin{figure} \\begin{center} \\includegraphics[width=8cm,height=6cm]{fig26.ps} \\caption{ALFOSC PSF obtained from stars in our frames. The PSF is the average of three stars, where one saturated star was used to obtain the far wings of the profile beyond $\\sim12$ arcsec. The individual star profiles were normalized to the total flux inside a $5$ arcsec radius before averaging.}.\\protect\\label{psffig} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "1208/1208.1768_arXiv.txt": {
        "abstract": "The variable magnetic field of the solar photosphere exhibits periodic reversals as a result of dynamo activity occurring within the solar interior. We decompose the surface field as observed by both the Wilcox Solar Observatory and the Michelson Doppler Imager into its harmonic constituents, and present the time evolution of the mode coefficients for the past three sunspot cycles.  The interplay between the various modes is then interpreted from the perspective of general dynamo theory, where the coupling between the primary and secondary families of modes is found to correlate with large-scale polarity reversals for many examples of cyclic dynamos.  Mean-field dynamos based on the solar parameter regime are then used to explore how such couplings may result in the various long-term trends in the surface magnetic field observed to occur in the solar case. ",
        "introduction": "\\label{sec:intro}  The Sun is a dynamic star that possesses quasi-regular cycles of magnetic activity having a mean period of about 22~yr.  This period varies from cycle to cycle, and over the past several centuries has ranged from 18 to 25 years \\citep{wei1990,bee1998,uso2007}, as for example illustrated by the unusual but not unprecedented length of the most recently completed sunspot cycle~23. During each sunspot cycle (comprising half of a magnetic cycle), the Sun emerges sunspot groups and active regions onto the photosphere, with such features possessing characteristic latitudes, polarity, and tilt angles.  As with the period, the numbers and emergence frequencies of active regions is observed to vary from cycle to cycle.  At activity minima when few active regions are present, the surface magnetic field is characterized by the presence of two polar caps, i.e., largely unipolar patches of magnetic flux dispersed across both polar regions with the northern and southern caps possessing opposite polarities.  Reversals of this large-scale dipole represented by the polar-cap flux occur during each sunspot cycle, allowing the subsequent sunspot cycle to begin in the opposite configuration.  After two sunspot cycles, and thus after undergoing two polarity reversals, the photospheric field will have returned to its starting configuration so as to complete a full activity cycle.  In response to the photospheric flux associated with various features, such as active regions and their decay products, the coronal magnetic field possesses structures having a broad spectrum of sizes.  These structures are both evident in observations of coronal loops, as found in narrow-band extreme ultraviolet or soft X-ray imagery, and reproduced in models of the coronal magnetic field (e.g., \\citealt{sch2003}).  In both venues, the coronal magnetic field geometry is seen to contain a rich and complex geometry. Dynamical events originating from the corona, such as eruptive flares and coronal mass ejections, are likely powered by energy released by a reconfiguration of the coronal magnetic field, which in turn is responding to changes and evolution of photospheric fields.  Precise measurements of the time-history of photospheric magnetic field and the ability to determine the projection of this field into its constituent multipole components are helpful in investigating the physical processes thought to be responsible for dynamo activity \\citep{bul1954,mof1978}.  In cool stars similar to the Sun, the dynamo is presumed to be a consequence of the nonlinear interactions between convection, rotation, and large scale flows, leading to the generation and maintenance against Ohmic diffusion of magnetic field of various temporal and spatial scales \\citep{wei1987,cat1999,oss2003,bru2004,voe2007,cha2010,rei2012}.  In particular, the dependence of dynamo activity upon rotation appears to be well established \\citep{noy1984,saa1999,piz2003,boh2007,rei2009}.  However, many details of the understanding of why many cool-star dynamos excite waves of dynamo activity having a regular period, specifically 22~yr in the case of the Sun, remain unclear.  To investigate this question, it is useful to explore the behavior and evolution of the lowest-degree (i.e., largest-scale) multipoles, their amplitudes and phases, and their correlations with the solar photospheric magnetic field.  Many earlier studies (e.g., \\citealt{lev1977b,hoe1984,gok1992a,gok1992b}) have illustrated how power in these modes ebb and flow as a function of the activity level.  In particular, J.~Stenflo and collaborators have performed thorough spectral analyses on the temporal evolution of the various spherical harmonic modes.  \\citet{ste1986}, \\citet{ste1987}, and \\citet{ste1988}, and more recently \\citet{ste1994} and \\citet{kna2005}, base their analysis on Mt.~Wilson and Kitt Peak magnetic data spanning the past few sunspot cycles.  As one would expect, they find that most of the power is contained in temporal modes having a period of about 22~yr, and especially in spherical harmonics that are equatorially antisymmetric, such as the axial dipole and octupole.  However, they find signatures of the activity cycle are present in all axisymmetric harmonics, as significant power is present at temporal frequencies at or near integer multiples of the fundamental frequency of 1.44~nHz [equivalent to (22~yr)$^{-1}$].  In the current study, we focus on the coupling between spherical harmonic modes, and what such coupling may indicate about the operation of the interior dynamo.  In particular, reversals of the axial dipole mode may be viewed as a result of continuous interactions between the poloidal and toroidal components of the interior magnetic field, i.e., the so-called dynamo loop.  Currently, one type of solar dynamo model that successfully reproduces many observed behaviors is the flux-transport Babcock-Leighton type (e.g., \\citealt{cho1995,dik2004,jou2007,yea2008}).  A key ingredient in producing realistic activity cycles using this type of model is found to be the amplitude and profile of the meridional flow \\citep{jou2007,kar2010,nan2011,dik2011}, which result in field reversals progressing via the poleward advection across the surface of trailing-polarity flux from emergent bipolar regions.  During the rising phase of each sunspot cycle, polar cap flux left over from the previous cycle is canceled, after which new polar caps having the opposite magnetic polarity form \\citep{wany1989,ben2004,das2010}.  Helioseismic analyses of solar oscillations have provided measurements and inferences of key dynamo components, such as the internal rotation profile and the near-surface meridional circulation \\citep{tho2003,bas2010}. Complementing precise observations of the solar magnetic cycle properties, these helioseismic inversions represent additional strong constraints on theoretical solar dynamo models.  Successful solar dynamo models strive to reproduce as many empirical features of solar magnetic activity as possible, including not only cycle periods, but also parity, phase relation between poloidal and toroidal components, and the phase relation between the dipole and higher-degree harmonic modes.  Interestingly, a recent analysis of geomagnetic records has indicated that the interplay between low-degree harmonic modes during polarity reversals is one way to characterize both reversals of the geomagnetic dynamo (which have a mean period of about 300,000~yr) as well as excursions, where the dipole axis temporarily moves equatorward and thus away from its usual position of being approximately aligned with the rotation axis, followed by a return to its original position without having crossed the equator (see \\citealt{hul2010} for a recent review on Earth's magnetic field).  In particular, these studies have shown that, during periods of geomagnetic reversals, the quadrupolar component of the geomagnetic field is stronger than the dipolar component, while during an excursion (which can be thought of as a failed reversal), the dipole remains dominant \\citep{ami2010,leo2007,leo2009}.  One may thus ask: Is a similar behavior observed for the solar magnetic field?  In an attempt to address this question, we have performed a systematic study of the temporal evolution of the solar photospheric field by determining the spherical harmonic coefficients for the photospheric magnetic field throughout the past three sunspot cycles, focusing on low-degree modes and the relative amplitude of dipolar and quadrupolar components.  Following the classification of \\citet{mcf1991}, we have made the distinction between primary and secondary families of harmonic modes, a classification scheme that takes into account the symmetry and parity of the spherical harmonic functions (see \\citealt{gub1993} for a detailed discussion on symmetry and dynamo).  While we recognize that the solar dynamo operates in a more turbulent parameter regime than the geodynamo, and is more regular in its reversals, the presence of grand minima (such as the Maunder Minimum) in the historical record indicates that the solar dynamo can switch to a more intermittent state on longer-term, secular time scales.  In fact in the late stages of the Maunder Minimum, the solar dynamo was apparently asymmetric, with the southern hemisphere possessing more activity than the north \\citep{rib1993} for several decades, a magnetic configuration that may have been achieved by having dipolar and quadrupolar modes of similar amplitude \\citep{tob1997,gal2009}. Additionally, recent spectopolarimetric observations of solar-like stars now provide sufficient resolution to characterize the magnetic field geometry in terms of its multipolar decomposition \\citep{pet2008b}.  Furthermore, the analysis of reduced dynamical systems developed over the last 20~yr describing the geodynamo and solar dynamo have emphasized the importance of the nonlinear coupling between dipolar and quadrupolar components \\citep{kno1996,wei2000,pet2009}.  This article is organized in the following manner.  In \\S\\ref{sec:observations}, we describe the data sets and the data analysis methods used to perform the spherical harmonic analysis, followed in \\S\\ref{sec:expansion} with an explanation of the temporal evolution of the various harmonic modes, the magnetic energy spectra, and the decomposition in terms of primary and secondary families.  We interpret in \\S\\ref{sec:discussions} our results from a dynamical systems perspective and illustrate some of these concepts using mean-field dynamo models.  Concluding remarks are presented in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Cycles of magnetic activity in many astrophysical bodies, including the Sun, Earth, and other stars, are thought to be excited by nonlinear interactions occurring in their interiors.  Yet in some cases, such as the Sun, the cycles have approximately regular periods and in others, such as the Earth, there is no apparent periodicity.  Dynamo theory indicates such a range of behaviors is expected, and whether the cycles are regular depends on magnetohydrodynamic parameters that characterize the system, including fluid and magnetic Reynolds and Rayleigh numbers.  As a consequence, the large-scale appearance of the magnetic field may provide clues toward the type of dynamo that may be operating.  In this article, long-term measurements of the solar photospheric magnetic field are utilized to characterize the waves of dynamo activity that exist within the interior of the Sun.  Synoptic maps from WSO (dating back to 1976) and MDI (spanning 1995--2010) are used to determine the spherical harmonic coefficients of the surface magnetic field for the past three sunspot cycles.  We focus on the apparent interactions between various low-order modes throughout the past three sunspot cycles, and interpret these trends in the context of dynamo theory.  The multipolar expansions of the solar field as deduced from WSO and MDI data indicate that the axial and equatorial dipole modes are out of phase.  During activity minima, the dipole component of the solar field is generally aligned with the axis of solar rotation, while the quadrupole component is much weaker.  During activity maxima, the dipole reverses its polarity with respect to the rotation axis, and throughout the reversal process there is more energy in quadrupolar modes than in dipole modes.  During the past three cycles, these reversals have taken place over a time interval of about 2~yr to 3~yr on average.  More indirect measures of solar activity, such as the sunspot number and proxies of the heliospheric field, seem to indicate that such regular activity cycles have persisted for at least hundreds of years with a period of approximately 11~yr.  The most recently completed solar cycle (Cycle~23) lasted for about 13~yr and while unusual, is not unprecedented.  We note in passing that such modulations of the solar dynamo may be interpreted as a type of nonlinear interaction between the turbulent alpha effect and the field and/or flows \\citep{tob2002}.  The harmonic modes can also be grouped into primary and secondary families, a distinction that depends on the north-south symmetry of the various modes. For example, the axial dipole harmonic is antisymmetric and is a member of the primary family.  Alternatively, the equatorial dipole and axial quadrupole modes are both symmetric with respect to the equator and thus are grouped together in the secondary family.  When the evolution of the mode coefficients are analyzed in this way, we find that there is a trend for members of the same family to possess the same phasing, suggesting that modes in the same family of modes are either excited together and/or are more coupled when compared with modes of different families.  This coupling is noticeable during reversals of the solar dipole, as less energy is present in primary-family modes than in secondary-family modes during these intervals.  The historical record indicates that the geodynamo also undergoes reversals of its dipole axis (with respect to the rotation axis), but these reversals occur much more irregularly than in the solar case.  Additionally, the dipole axis of the terrestrial magnetic field occasionally makes excursions away from the axis of rotation of the Earth, only to later return without actually reversing.  An examination of the large-scale harmonic modes of the geomagnetic field during these intervals indicates that the energy contained in secondary-family modes was significantly smaller during excursions than during reversals.  A strong quadrupole during geodynamo reversals is in line with the solar behavior; there is no parallel with excursions as excursions in the solar case have not been observed.  Analogous behavior is observed to occur in the VKS laboratory dynamo with respect to the relative strengths of the primary and secondary families.  We also examined the coupling of the mode families using a BL mean-field dynamo model computed using the STELEM code.  Because of the symmetries in the magnetic induction and in the assumed profiles of the large-scale flow fields and BL source term, we find that the standard mean-field solar dynamo model results in a state containing largely members of the primary family.  This is a result of the dipole (a primary-family mode) being more unstable to dynamo action than the quadrupole.  With a modest amount of asymmetry, implemented here either in the meridional flow profile or in the BL source term, we find from the models that both the primary and secondary families can coexist in the same model and in the same proportions as in the solar dynamo.  This can lead to a small lag between the northern and southern hemispheres as is actually observed on the Sun \\citep{dik2007}."
    },
    "1208/1208.3906_arXiv.txt": {
        "abstract": "The inspiral of a stellar mass ($1 - 100\\,M_\\odot$) compact body into a massive ($10^5 - 10^7\\,M_\\odot$) black hole has been a focus of much effort, both for the promise of such systems as astrophysical sources of gravitational waves, and because they are a clean limit of the general relativistic two-body problem.  Our understanding of this problem has advanced significantly in recent years, with much progress in modeling the ``self force'' arising from the small body's interaction with its own spacetime deformation.  Recent work has shown that this self interaction is especially interesting when the frequencies associated with the orbit's $\\theta$ and $r$ motions are in an integer ratio: $\\Omega_\\theta/\\Omega_r = \\beta_\\theta/\\beta_r$, with $\\beta_\\theta$ and $\\beta_r$ both integers.  In this paper, we show that key aspects of the self interaction for such ``resonant'' orbits can be understood with a relatively simple Teukolsky-equation-based calculation of gravitational-wave fluxes.  We show that fluxes from resonant orbits depend on the relative phase of radial and angular motions.  The purpose of this paper is to illustrate in simple terms how this phase dependence arises using tools that are good for strong-field orbits, and to present a first study of how strongly the fluxes vary as a function of this phase and other orbital parameters.  Future work will use the full dissipative self force to examine resonant and near resonant strong-field effects in greater depth, which will be needed to characterize how a binary evolves through orbital resonances. ",
        "introduction": "\\label{sec:intro}  \\subsection{The self-force driven evolution of binaries:\\\\ A very brief synopsis} \\label{sec:synopsis}  Our understanding of the two-body problem in general relativity has advanced substantially in the past decade.  Besides the celebrated breakthroughs in numerical relativity {\\cite{pretorius05, campanelli06, baker06}} which have opened the field of binary phenomenology in general relativity, there has been great progress in understanding the extreme mass-ratio limit of this problem, when one member of the binary is much smaller than the other.  This limit is of great interest in describing astrophysical extreme mass-ratio binaries (a particularly interesting source for space-based gravitational-wave measurements) {\\cite{emri}}, and as a limiting form of the more generic two-body problem {\\cite{lousto,sperhake}}.  Most efforts to model extreme mass-ratio binaries have focused on the computation of {\\it self forces} (see Ref.\\ {\\cite{barack}} for a recent comprehensive review).  Consider a small body orbiting a black hole.  At zeroth order in the small body's mass, its motion is described as a geodesic of the black hole spacetime.  At first order in this mass, the black hole's spacetime is slightly deformed.  This deformation changes the trajectory that the small body follows, pushing it away from the background spacetime's geodesic.  It is useful to regard the change to the trajectory as arising from a self force which modifies the geodesic equations typically used to describe black hole orbits.  Conceptually, it is useful to split the self force into two pieces: A time-symmetric {\\it conservative} piece, and a time-asymmetric {\\it dissipative} piece.  On average, the impact of the conservative contribution is to shift orbital frequencies away from their geodesic values.  The dissipative self force is equivalent, on average, to a slow evolution of the otherwise conserved constants (e.g., the orbital energy and angular momentum) which characterize geodesic orbits.  It makes the largest contribution to an orbit's phase evolution.  The conservative piece makes a smaller (though still significant) contribution which accumulates secularly over many orbits {\\cite{ppn,pp}}.  Recent work by Flanagan and Hinderer {\\cite{FH}} (hereafter FH) using a post-Newtonian (pN) approximation to the self force together with fully relativistic orbital dynamics has shown that a small body's self interaction becomes particularly important near {\\it resonances}.  The background geodesic motion can be characterized by three orbital frequencies with respect to Boyer-Lindquist time: A radial frequency $\\Omega_r$, a polar frequency $\\Omega_\\theta$, and an axial frequency $\\Omega_\\phi$.  In the weak-field (large separation) limit, these three frequencies asymptote to the Newtonian Kepler frequency.  In the strong field, these frequencies can differ significantly, with $\\Omega_r$ always the smallest frequency (the relative magnitude of $\\Omega_\\theta$ and $|\\Omega_\\phi|$ depends on the sign of the orbit's axial angular momentum). Resonant orbits are ones for which the radial and angular motions become commensurate: $\\Omega_\\theta/\\Omega_r = \\beta_\\theta/\\beta_r$, where $\\beta_\\theta$ and $\\beta_r$ are small integers with no common factors.  On such orbits, components of the self interaction which normally ``average away'' when examined over a full orbital period instead combine coherently, substantially changing their impact on the system's evolution.  For the purpose of our background discussion, it is useful to include more details from FH's analysis of how resonant effects arise. Consider a body of mass $\\mu$ moving on a bound trajectory near a Kerr black hole of mass $M$, with $\\mu \\ll M$.  FH note that one can describe the motion of this body using action-angle variables and correctly accounting for how the integrals which parameterize geodesic orbits evolve due to the self force.  Writing the angle variables $q_\\alpha = (q_t,q_r,q_\\theta,q_\\phi)$ (which describe motions in the $t$, $r$, $\\theta$, and $\\phi$ directions of Boyer-Lindquist coordinates), and writing the integrals associated with geodesic motion $J_i = (E, L_z, Q)$ (with $E$ the energy, $L_z$ the axial angular momentum, and $Q$ the Carter constant), the equations of motion describing the system are {\\cite{FHtt}} \\begin{eqnarray} \\frac{dq_\\alpha}{d\\tau} &=& \\omega_\\alpha({\\bf J}) + \\epsilon g^{(1)}_\\alpha(q_r,q_\\theta,{\\bf J}) + O(\\epsilon^2)\\;, \\label{eq:eom1}\\\\ \\frac{dJ_i}{d\\tau} &=& \\epsilon G^{(1)}_i(q_r,q_\\theta,{\\bf J}) + O(\\epsilon^2)\\;. \\label{eq:eom2} \\end{eqnarray} The time parameter $\\tau$ is proper time along the orbit; the parameter $\\epsilon = \\mu/M$, the system's mass ratio.  The $\\omega_{r,\\theta,\\phi}$ are fundamental frequencies with respect to proper time associated with bound Kerr geodesic orbits.  The forcing functions $g^{(1)}_\\alpha$ and $G^{(1)}_i$ arise from the first-order self force.  FH also include discussion of second-order forcing functions, which we do not need for this synopsis; see Ref.\\ {\\cite{FH}} for further discussion.  At order $\\epsilon^0$, Eqs.\\ (\\ref{eq:eom1}) and (\\ref{eq:eom2}) simply describe geodesics of Kerr black holes: The integrals of the motion are constant, and each angle variable evolves according to its associated frequency.  The leading adiabatic dissipative correction to this motion can be found by dropping the forcing term $g^{(1)}_\\alpha$ and replacing $G^{(1)}_i$ by $\\langle G^{(1)}_i\\rangle$, the average of this forcing term over the 2-torus parameterized by $q_\\theta$ and $q_r$ {\\cite{FHtt}}.  To compute this torus-averaged self force, it is sufficient to use the radiative approximation {\\cite{mino03,FHtt,pp}}, which includes only the radiative contributions to the self interaction and neglects conservative contributions.  For generic (non-resonant) orbits, this torus average coincides with an infinite time average, and the averaged quantities $\\langle G^{(1)}_i \\rangle$ are just the time-averaged fluxes of energy, angular momentum and Carter constant.  In recent years such time-averaged fluxes have been computed numerically using the frequency domain Teukolsky equation {\\cite{dh06,fujita,sago}}.  These fluxes can be used to compute leading-order, adiabatic inspirals.  The conservative contributions influence the motion only beyond the leading adiabatic order {\\cite{mino03,FHtt}}.  \\subsection{Resonant effects} \\label{sec:res_intro}  Now consider going beyond the leading adiabatic order.  Important post-adiabatic effects can be found by continuing to neglect $g^{(1)}_\\alpha$, but now integrating Eq.\\ (\\ref{eq:eom2}) using $G^{(1)}_i$ rather than its averaged variant.  FH show that for ``most'' orbits, $G^{(1)}_i$ is given by $\\langle G^{(1)}_i\\rangle$ plus a rapidly oscillating contribution.  Over the timescales associated with inspiral, this rapidly oscillating piece averages away and has little effect.  The effect of the forcing term $G^{(1)}_i$ is dominated by $\\langle G^{(1)}_i\\rangle$ for all non-resonant orbits.  For resonant orbits, this averaging fails: contributions beyond $\\langle G^{(1)}_i\\rangle$ are {\\it not} rapidly oscillating, and can significantly modify how the integrals of motion evolve during an inspiral.  A given binary is very likely to evolve through several low-order resonances en route to the final merger of the smaller body with the large black hole {\\cite{rh_inprep}}.  A complete quantitative understanding of these resonant effects will thus be quite important for making accurate inspiral models.  Prior to FH's analysis, several other papers argued that such resonances may play an important role in the radiative evolution of binary systems {\\cite{mino05,tanaka06}} (albeit without quantifying the detailed impact they can have), or else because of other effects which resonances have on the evolution of a dynamical system {\\cite{acl10}}.  Orbits in which $\\Omega_\\theta/\\Omega_r$ take on a small-integer ratio have been studied in great detail by Grossman, Levin, and Perez-Giz {\\cite{glpgI}}, who called them ``periodic'' orbits and provided a fairly simple scheme for classifying their features.  Following Ref.\\ {\\cite{FH}} (as well as more recent work by Grossman, Levin, and Perez-Giz {\\cite{glpgII}}), we will call them ``resonant'' orbits, reflecting the fact that our main interest is in understanding how their periodic structure impacts the self interaction.  Grossman, Levin and Perez-Giz have more recently argued for the utility of using resonant orbits as sample points in numerical computations of leading order, adiabatic inspirals: evaluating fluxes at resonant orbits may enable a speedup of flux computations {\\cite{glpgII}}, more efficiently covering the parameter space of generic orbits.  Although their goals are rather different from ours here, many of their techniques and results substantially overlap with ours (modulo minor differences in notation).  We highlight the overlap at appropriate points in this paper.  As a binary evolves through a resonance, its self interaction and thus its evolution are modified compared to what we would expect if the resonance were not taken into account.  The details of how the self interaction is modified depend on the relative phase of the radial and angular motions as the orbit passes through resonance.  Because of this, {\\it resonances enhance the dependence of a binary's orbital evolution on initial conditions.}  Let the phase variable $\\chi_0$ define the value of the orbit's $\\theta$ angle at the moment it reaches periapsis (see Sec.\\ {\\ref{sec:generalgeod}} for more details).  On resonance, two orbits which have the same energy $E$, the same axial angular momentum $L_z$, and the same Carter constant $Q$ will evolve differently if they have different values of $\\chi_0$.  FH estimate {\\cite{FH}} that the shift to the orbital phase induced by these resonances can be several tens to $\\sim 10^2$ radians for mass ratios $\\sim 10^{-6}$ (as compared to an analysis which neglects the resonances).  That there is such a large shift, and that this shift may depend on initial conditions, is potentially worrisome. Resonances could significantly complicate our ability to construct models for measuring the waves from extreme mass-ratio inspirals.  On the other hand, the detailed behavior of a system as it evolves through resonances may offer an opportunity to study an interesting aspect of strong-field gravity, providing a new handle for strong-gravity phenomenology.  Analytic studies of the effect of the passage through resonance can be found in Refs.\\ \\cite{gyb11,fh13}.  \\subsection{Our analysis} \\label{sec:our_paper}  The ``several tens to $\\sim 10^2$ radians'' estimate by FH is based on applying pN self force estimates to strong-field orbits, a regime where pN approximations are generally inaccurate.  It is thus of great interest to estimate the impact of orbital resonances using strong-field methods.  The purpose of this paper is to take a first step in this direction.  Our goal is to generalize our computational techniques in order to treat resonances correctly.  A key point is that the flux-balancing technique which can be used to approximate inspiral (as described in the final paragraph of Sec.\\ {\\ref{sec:synopsis}}) is based on the adiabatic approximation.  This approximation temporarily breaks down during a resonance.  Therefore, to treat resonances, one must use the orbital equations of motion (\\ref{eq:eom1}) -- (\\ref{eq:eom2}) to track the evolution of all the orbital degrees of freedom on short timescales.  Flux balancing instead just tracks the evolution of the conserved quantities $E$, $L_z$ and $Q$ on long timescales.  In addition one must use the full, oscillatory self-force driving term $G^{(1)}_i$, and not just its averaged version.  As is well known, computation of the full self force is extremely difficult, largely because it requires regularization of the self field \\cite{barack}.  Fortunately, only the dissipative piece of the self force should contribute to leading order resonance effects.  As argued in FH, there is some evidence suggesting that geodesic motion perturbed by the conservative piece of the self force is an integrable dynamical system, and resonances do not occur in such systems.  Thus, if the integrability conjecture of FH is true, only the dissipative self force needs to be computed.  This constitutes a great simplification, since the well-known difficulties of self-force computations apply only to the conservative piece; the dissipative piece is relatively straightforward to compute.  Techniques for doing so with scalar fields were presented in Ref.\\ {\\cite{dfh05}}, and generalizing to the gravitational dissipative self force is not terribly difficult {\\cite{mino05,tanaka06}}.  While these references focused on the {\\it averaged} self force, it is straightforward to generalize the analysis to obtain the full dissipative self force.  It is thus feasible to perform numerical compututations of orbital evolutions through resonances using the full dissipative self force, without any orbit averaging.  Our eventual goal is to extend our black hole perturbation theory codes to do just this, and to evaluate how the dissipative self force behaves as a system evolves through resonance.  Work in this vein is in progress , and will be presented in future work {\\cite{fhhr_inprep}}.  In this paper, we take a first step in this direction.  We focus here on computation of time-averaged fluxes of the integrals of the motion, and in particular on how these quantities differ between resonant and non-resonant orbits.  These quantities correspond to the fluxes that one would measure at infinity (and at the black hole horizon) if one turned off radation reaction effects; upon averaging over long times, they are equal to the rate at which the dissipative self force evolves these constants.  We emphasize that these quantities are not sufficient to allow computation of orbital evolutions.  However, they provide insight into the characteristic features of the radiation emitted by resonant orbits.  We find that fluxes from resonant orbits generically differ from those from nearby, non-resonant orbits\\footnote{Thus the fluxes change discontinuously as one varies the orbital parameters.  This is certainly unphysical, but arises because we compute infinite time averages of fluxes from geodesic orbits.  If one considers the fluxes from the true inspiraling motion, and averages over a timescale intermediate between the orbital timescale and the radiation reaction timescale, the time-averaged fluxes would vary smoothly with time, with order unity changes in the vicinity of resonances.  This point is discussed further in Appendix {\\ref{app:qdot}}.}, and in addition vary depending on the relative phase of the radial and angular motions.  The magnitude of this variation is closely related to the ``kick'' that is imparted to the orbit's constants as it evolves through a resonance (cf.\\ Fig.\\ 1 of FH).  As such, characterizing on-resonance fluxes is a useful and natural first step in the process of modifying existing flux-based codes to compute the full dissipative self force.  We explore numerically the magnitude of the difference between the resonant and non-resonant cases, and the dependence on the orbital phase.  For specific modes, the fluxes can vary by large factors (although variations of order unity are more typical).  For the net fluxes obtained by summing over all modes, variations are typically of order a percent or less.  \\subsection{Outline of this paper} \\label{sec:outline}  We begin this paper by briefly reviewing the behavior of Kerr geodesic orbits in Sec.\\ {\\ref{sec:geodesics}}.  Much of this material has been presented elsewhere, so we leave out most details, pointing the reader to appropriate references.  Our main focus is to describe how to find and characterize resonant orbits.  We then describe how to compute radiation from Kerr orbits in Sec.\\ {\\ref{sec:radiate}}.  We first briefly review the Teukolsky-equation-based formalism we use (Secs.\\ {\\ref{sec:fdteuk}} -- {\\ref{sec:fluxes}}), and then describe how key details are modified by orbital resonances in Secs.\\ {\\ref{sec:resI}} and {\\ref{sec:resII}}.  We describe two complementary approaches to computing fluxes on resonance.  Although formally equivalent (as we prove in Appendix {\\ref{app:equivalence}}), their implementation is quite different.  Having both methods at hand proved useful to us in our numerical study.  One aspect of the on-resonance computation (the evolution of Carter's constant $Q$) is sufficiently complicated that all details of this calculation are given in Appendix {\\ref{app:qdot}}.  Our analytic results for fluxes of energy and angular momentum on resonance agree with those obtained by Grossman, Levin and Perez-Giz (compare especially Secs.\\ IIID--E and Appendices B5, B6, and C in Ref.\\ {\\cite{glpgII}} with our discussion here, and with our Appendix {\\ref{app:equivalence}}).  Our result for the resonant rate of change of the Carter constant appears to be new.  Our numerical results are given in Sec.\\ {\\ref{sec:results}}.  We begin by examining how fluxes from individual modes (i.e., harmonics of the orbital frequencies) behave as a function of the offset phase of the radial and angular motions, which we denote $\\chi_0$.  We show that the amplitude of a given mode, and hence the rates of change of conserved quantities associated with that mode, can vary significantly with $\\chi_0$.  For example, the flux of energy from an orbit can vary by factors of order unity as $\\chi_0$ varies from $0$ to $2\\pi$.  The rate of change of the orbit's Carter constant can even change sign as $\\chi_0$ varies.  The total flux from a given orbit is given, however, by adding fluxes from many modes.  When many modes are combined, much of the variation washes away; we find variations of a fraction of a percent in most quantities after summation.  The amount of this residual variation seems to depend most strongly upon the geometry of the orbit's $(r,\\theta)$ motion on resonance, in particular the topology of the trace in the $(r,\\theta)$ plane.  Orbits whose motion in $(r,\\theta)$ have a simple topology with few trajectory crossings in the plane (e.g., the $\\Omega_\\theta/\\Omega_r = 3/2$ resonance) tend to have relatively large variation in the integrals of motion; orbits whose motion has a more complicated topology with many trajectory crossings show much less variation (e.g., the $\\Omega_\\theta/\\Omega_r = 4/3$ resonance).  We argue that this can be explained in terms of how the orbital motion tends (or fails) to average away variations in the source-term to the Teukolsky equation.  As emphasized in Sec.\\ {\\ref{sec:our_paper}}, understanding these fluxes exactly on resonance is only the first step in building a complete strong-field understanding of how resonances impact inspirals.  In particular, these results do not provide enough information to specify how a system will evolve through a resonance. To go further, it will be necessary to examine how dissipation behaves as the system evolves toward and through an orbital resonance.  As mentioned above, this analysis is now beginning; we briefly outline the approach we are pursuing in Sec.\\ {\\ref{sec:conclude}}.  Throughout this paper, we use ``relativist's units,'' setting $G = 1 = c$. ",
        "conclusions": "\\label{sec:conclude}  In this analysis, using a Teukolsky-equation-based formalism good for exploring radiation produced by strong-field orbits, we have confirmed the picture that on resonance the gravitational-wave driven evolution of a binary can depend strongly on the relative phase of radial and angular motions.  A binary in which this relative phase has the value $\\pi/2$ as the system enters resonance may evolve quite differently from an otherwise identical system in which this phase is $3\\pi/2$ entering resonance.  A typical extreme mass-ratio binary can be expected to pass through several orbital resonances en route to its final coalescence.  That their evolution through each resonance depends strongly on an ``accidental'' phase parameter has the potential to complicate schemes for measuring gravitational waves from these binaries.  We find that the degree of variation depends strongly upon the topology of the orbital trajectory in the $(r,\\theta)$ plane\\footnote{Strictly speaking, it is a trajectory's geometry that matters, particularly how close the orbit comes to all accessible points in the $(r,\\theta)$ plane.  However, its geometry is strongly correlated to its topology, which is an invariant property of a resonant orbit's frequencies \\cite{glpgI}.  As such, the topology is a valuable way to characterize this aspect of its resonant behavior.}.  Of the cases we have studied in detail, the orbital plane trajectory of resonances like $\\Omega_\\theta/\\Omega_r = 3/2$ have a simple topology.  This trajectory does not cross itself very often, and does not come close to many points in the plane.  Such resonances do not effectively average out the behavior of the source to the wave equation.  As such, if the source varies significantly over an orbit, there can be a strong residue of this variation in the associated radiation.  By contrast, the trajectory of resonances like $\\Omega_\\theta/\\Omega_r = 4/3$ has a more complicated topology, crossing itself many times, and more completely ``covering'' the plane.  In these cases, the orbit comes ``close to'' many of the allowed points in the $(r,\\theta)$ plane, which quite effectively averages out the source's behavior.  Although instructive and a nice validation of our ability to examine resonances, these results are not enough to truly assess the importance that resonances have in a strong field analysis.  We must be able to analyze a system as it evolves through a resonance, and thereby integrate the full ``kicks'' in the integrals of motion $E$, $L_z$, and $Q$ imparted to the system as it passes through resonance. A first step in this direction has been taken by van de Meent {\\cite{vdM}}, who examines the likelihood that resonances can ``trap'' an orbit, leading to long-lived resonant waves.  Part of van de Meent's analysis is adescription of the system's evolution as motion through a one-dimensional effective potential.  This approach is likely to be useful for more general analysis of resonant evolution.  For our planned work, we have begun expanding our Teukolsky code to compute, in the frequency domain, the instantaneous components of the dissipative or radiative piece of the self force.  Our formulation is based in part on the discussion of Refs.\\ {\\cite{mino05,tanaka06,dfh05}}, but generalized to compute the full dissipative self force rather than its torus average\\footnote{One might be concerned about gauge ambiguities associated with the gravitational self force.  As shown by Mino {\\cite{mino03}}, these ambiguities disappear when one averages the self force's effects over an infinite time.  In a two-timescale expansion {\\cite{FHtt}}, such ambiguities remain, but are suppressed by the ratio of the timescales.}.  This will allow us to study how a real inspiral is affected as we evolve through each resonance using results that are good deep in the strong field.  The results shown in this paper are a first step toward this, demonstrating that our strong-field toolkit can be used to study resonant effects."
    },
    "1208/1208.0714_arXiv.txt": {
        "abstract": "We present a new measurement of the $\\alpha$-spectroscopic factor ($S_\\alpha$) and the asymptotic normalization coefficient (ANC) for the 6.356 MeV 1/2$^+$ subthreshold state of $^{17}$O through the $^{13}$C($^{11}$B,\\,$^{7}$Li)$^{17}$O transfer reaction and we determine the $\\alpha$-width of this state. This is believed to have a strong effect on the rate of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction, the main neutron source for {\\it slow} neutron captures (the $s$-process) in asymptotic giant branch (AGB) stars. Based on the new width we derive the astrophysical S-factor and the stellar rate of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction. At a temperature of 100 MK our rate is roughly two times larger than that by \\citet{cau88} and two times smaller than that recommended by the NACRE compilation. We use the new rate and different rates available in the literature as input in simulations of AGB stars to study their influence on the abundances of selected $s$-process elements and isotopic ratios. There are no changes in the final results using the different rates for the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction when the $^{13}$C burns completely in radiative conditions. When the $^{13}$C burns in convective conditions, as in stars of initial mass lower than $\\sim$2 $M_\\sun$ and in post-AGB stars, some changes are to be expected, e.g., of up to 25\\% for Pb in our models. These variations will have to be carefully analyzed when more accurate stellar mixing models and more precise observational constraints are available. ",
        "introduction": "Approximately half of the elements heavier than iron in the universe are produced via a series of $slow$ neutron capture reactions and competing $\\beta$-decays (the $s$-process). During the $s$-process, the neutron number density is relatively low, of the order of 10$^7$ n cm$^{-3}$. When the flux reaches an unstable nucleus, it typically decays rather than capture another neutron and the $s$-process proceeds via the isotopes around the valley of $\\beta$-stability \\citep[e.g.,][]{bur57}. The astrophysical sites of the $s$-process are core He and shell C burning in massive stars for the elements lighter than Sr \\citep{pig10}, and the ``He intershell'' of asymptotic giant branch (AGB) stars for the elements between Sr and Bi \\citep{gal98}. Stars with initial masses lower than roughly 9 $M_\\sun$ reach the AGB phase in the final phases of their evolution, when both H and He have been exhausted in the core leaving C and O in electron degenerate conditions. Production of nuclear energy occurs in the H and He shells, which are located between the core and the extended convective envelope and are separated by the thin He intershell layer. AGB stars experience thermal pulses (TPs) when the usually dormant He burning shell is suddenly activated. A large amount of energy is released, which drives convection in the He intershell. During TPs the star expands and cools and the H burning shell is inactive. While He burning turns from the convective to the radiative regime, and eventually switches off, the convective envelope can penetrate the underlying He intershell and carry to the surface the products of He burning, in particular carbon and the elements heavier than iron made by the $s$-process. This mixing process is known as the ``third dredge-up'' (TDU). After the TDU is ended, the star contracts and heats up again and H burning resumes until another TP occurs and the cycle is repeated. This sequence of events can occur from a few times to hundreds of times, depending on the stellar mass and the mass-loss rate. AGB stars suffer from very strong stellar winds, which erode the envelope roughly within a million years and shed the newly synthesized material mixed to the surface by the TDU into the interstellar medium. Eventually, the C-O degenerate core is left as a cooling white dwarf \\citep[see][for a review on AGB stars]{her05}.  According to the current standard model \\citep{gal98,bus99,gor00,lug03a,cri09a}, some protons must diffuse from the convective envelope into the He intershell at the end of each TDU in order to produce enough $^{13}$C to account for the observed abundances of the $s$-process elements at the surface of AGB stars \\citep[see also][]{bus01}. A thin layer is then produced, known as the $^{13}$C ``pocket'', which is rich in $^{13}$C made via $^{12}$C($p$,\\,$\\gamma$)$^{13}$N($\\beta^+\\nu$)$^{13}$C. When in this region the temperature reaches about $9 \\times 10^7$ K, the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction is activated and generates neutrons that trigger the $s$-process \\citep{hol88,gal88,kap90,kap11}.  Considerable effort has been devoted to the direct measurement of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O cross section \\citep{sek67,dav68,bai73,kel89,dro93,bru93,har05,hei08}. These measurements have been performed at energies down to 270 keV, whereas the Gamow window is at 190 $\\pm$ 40 keV, corresponding to a temperature of 100 MK. Since this energy is far below the Coulomb barrier, the reaction cross section is extremely small and direct measurement is sensitively limited by background signals and very difficult to perform in laboratories on the Earth's surface. While a measurement has been proposed at the underground laboratory of LUNA \\citep{cos09}, at present, the experimental cross sections have to be extrapolated below 270 keV. A microscopic cluster model analysis of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O and $^{13}$C($\\alpha$,\\,$\\alpha$) reactions by \\citet{des87} suggested that this extrapolation is critically affected by the 1/2$^+$ subthreshold resonance in $^{17}$O ($E_x$ = 6.356 MeV, just 3 keV below the $\\alpha$-decay threshold). The contribution from this resonance depends strongly on the $\\alpha$-width of the 1/2$^+$ state in $^{17}$O, which can be derived from the spectroscopic factor ($S_\\alpha$) or the asymptotic normalization coefficient (ANC) of $\\alpha$-cluster in this state.  The $S_\\alpha$ and the ANC can be determined from the angular distribution of the direct $\\alpha$-transfer reaction using distorted wave Born approximation (DWBA) or coupled reaction channels (CRC) analysis. Although three indirect measurements via the ($^{6}$Li,\\,$d$) or the ($^{7}$Li,\\,$t$) system have been performed by \\citet{kub03}, \\citet{joh06}, and \\citet{pel08} to study the $S_\\alpha$ or the ANC of the 1/2$^+$ state, there still exists a significant discrepancy of up to a factor of $\\sim$30 in the derived $S_\\alpha$ and ANC. Therefore, it is interesting to perform a new measurement of the $S_\\alpha$ and the ANC via an independent transfer reaction. In addition, it is necessary to understand the impact of the different resulting $^{13}$C($\\alpha$,\\,$n$)$^{16}$O rates on the $s$-process nucleosynthesis in AGB stars.  In this paper we determine a new stellar rate of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction and incorporate it in calculations of the $s$-process nucleosynthesis in AGB stars. First, we measure the angular distribution of the $^{13}$C($^{11}$B,\\,$^{7}$Li)$^{17}$O reaction to determine the $S_\\alpha$ and the ANC for the 1/2$^+$ state in $^{17}$O. Using this experimental ANC we derive the $\\alpha$-width for the 1/2$^+$ subthreshold resonance, which is currently the most uncertain parameter for determining the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O rate. Finally, we use the new rate and different rates available in the literature as input for simulations of AGB stars to study their influence on the abundance of some selected $s$-process elements and isotopic ratios. ",
        "conclusions": "We determined the stellar rate of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction and incorporated the new reaction rate in calculations of the $s$-process nucleosynthesis in AGB stars. The $S_\\alpha$ and ANC for the 6.356 MeV 1/2$^+$ subthreshold state in $^{17}$O were obtained from the measurement of the $^{13}$C($^{11}$B,\\,$^{7}$Li)$^{17}$O angular distribution. This provided an independent examination to shed some light on the existing discrepancies in the $S_\\alpha$ and ANC values derived from different authors. Based on the measured ANC, we extracted the $\\alpha$-width of the 1/2$^+$ state in $^{17}$O, which is currently the most uncertain parameter for determining the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction rate. By using the present $\\alpha$-width and considering the properties of $^{17}$O states up to 8.342 MeV as well as their interferences, we derived the astrophysical S-factor and the stellar rate of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction. At a temperature of 100 MK the new rate is roughly two times larger than the CF88 value and two times smaller than that recommended by NACRE (see Fig. \\ref{fig5}). Verification of the present result using other independent techniques is desirable, e.g., the Trojan horse approach \\citep{spi99,muk08}, and the isospin symmetry approach based on a measurement of the 1/2$^+$ 6.560 MeV state in $^{17}$F \\citep{tim07}. In addition, an extension of the experimental data of $^{13}$C($\\alpha$,\\,$n$)$^{16}$O toward lower energies is highly desirable, which can probably only be performed in an underground laboratory, e.g., LUNA \\citep[see][]{cos09}.  We incorporated different $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction rates in calculations of the $s$-process nucleosynthesis in AGB stars and found that: (1) If $^{13}$C burns completely in radiative conditions during the interpulse phase (as for stars of initial mass greater than $\\sim$2 $M_\\sun$), there is no change in the final results. (2) If some $^{13}$C burns instead inside the convective TPs (for stars of initial mass lower than $\\sim$2 $M_\\sun$), we find changes of up to 25\\% in the $s$-process results, particularly for Pb. There are model uncertainties related to the result of point (2): a) when $^{13}$C burning in the TPs is due to incomplete burning of the $^{13}$C during the interpulse period, the exact stellar mass and metallicity range where incomplete burning of the $^{13}$C during the interpulse period occurs, as well as TP numbers, and the exact amount of $^{13}$C ingested, all depend sensitively on the temperature and density in the $^{13}$C pocket, on the interpulse period, and on the details of the inclusion of the $^{13}$C pocket. b) When $^{13}$C burning in the TPs is due to ingestion of protons directly inside the TP \\citep[as in low-mass and low-metallicity stars, as well as in post-AGB stars experiencing a late TP, see][]{her11}, the amount of $^{13}$C and neutrons produced strongly depends on the physical and numerical treatment of the mixing scheme adopted, which is at present uncertain.  Due to these model uncertainties together with the fact that the stellar observations have relatively large error bars, it is presently not possible to conclude if the new rate of the $^{13}$C($\\alpha$,\\,$n$)$^{16}$O reaction provides the best match to the available observational constraints. This may be however possible in the future, when development of recent 3D hydrodynamical models of the proton ingestion episodes \\citep{sta11} will allow a better understanding of neutron production and the $s$-process inside TPs to be compared to the composition of stellar observations and stardust grains."
    },
    "1208/1208.5900_arXiv.txt": {
        "abstract": "We present the discovery and extensive early-time observations of the Type Ic supernova (SN) PTF12gzk. Our light curves show a rise of 0.8\\,mag within 2.5\\,hr. Power-law fits [$f(t)\\propto(t-t_0)^n$] to these data constrain the explosion date to within one day. We cannot rule out a quadratic fireball model, but higher values of $n$ are possible as well for larger areas in the fit parameter space. Our bolometric light curve and a dense spectral sequence are used to estimate the physical parameters of the exploding star and of the explosion.  We show that the photometric evolution of PTF12gzk is slower than that of most SNe~Ic. The high ejecta expansion velocities we measure ($\\sim30,000$\\,km\\,s$^{-1}$ derived from line minima four days after explosion) are similar to the observed velocities of broad-lined SNe~Ic associated with gamma-ray bursts (GRBs) rather than to normal SN~Ic velocities. Yet, this SN does not show the persistent broad lines that are typical of broad-lined SNe~Ic. The host-galaxy characteristics are also consistent with GRB-SN hosts, and not with normal SN~Ic hosts. By comparison with the spectroscopically similar SN~2004aw, we suggest that the observed properties of PTF12gzk indicate an initial progenitor mass of 25--35\\,M$_\\odot$ and a large [(5--10) $\\times 10^{51}$\\,erg] kinetic energy, the later being close to the regime of GRB-SN properties. ",
        "introduction": "A core-collapse supernova (CCSN) occurs when a star having an initial mass $M\\geq8\\,{\\rm M}_\\odot$ ends its life in a catastrophic explosion. Observationally, CCSNe are divided into three groups based on their observed spectra: SNe~II show large amounts of hydrogen, SNe~Ib exhibit helium but little or no hydrogen, and SNe~Ic do not show significant amounts of hydrogen or helium \\citep[for a review, see][]{Filippenko1997}.  SNe~Ic are heterogeneous. Their luminosity, ejected mass, and kinetic energy span over an order of magnitude, from the subluminous SN 2004aw to the overluminous SN 1998bw \\citep{Mazzali2009,Drout2011}. The light-curve shapes of different events are also quite diverse. A subclass of SNe~Ic whose spectra are characterized by broad lines (Type Ic-BL; prototype SN 1998bw) is the only one for which clear evidence of an association with gamma-ray bursts (GRBs) exists (GRB-SNe; see Woosley \\& Bloom 2006 for a review). Superluminous SNe (SLSNe) of Type Ic are even more powerful \\citep[][and references therein]{GalYam2012}, but these probably result from a different physical mechanism. While SNe~Ic are common in the center of high-metallicity galaxies \\citep{Anderson2012}, SLSNe-I and broad-lined GRB-SNe tend to be found in dwarf hosts \\citep[e.g.,][]{Stanek2006,Modjaz2008,Arcavi2010}, giving untargeted sky surveys an advantage over targeted surveys in detecting these types of cosmic explosions.  The Palomar Transient Factory (PTF; Law et al. 2009; Rau et al. 2009) is a wide-field untargeted sky survey which explores the transient optical sky. It uses the PTF CFH12k camera mounted on the Palomar 48-inch telescope (P48).  PTF's short observing cadence and real-time capability \\citep[e.g.,][]{GalYam2011} enables the discovery and study of SNe at early stages of the explosion. In this {\\it Letter} we report the discovery and study of PTF12gzk, a peculiar SN~Ic in a dwarf star-forming galaxy located at redshift $z = 0.0137$ (distance 57.8\\,Mpc, distance modulus 33.8\\,mag, assuming H$_0$ = 71\\,km\\,s$^{-1}$\\,Mpc$^{-1}$). ",
        "conclusions": "PTF12gzk is a luminous SN~Ic, at the high end of the SN~Ic luminosity distribution \\citep{Drout2011}. It exhibits a slow rise of 18 days to its peak $r$-band magnitude, with $B$ peaking $\\sim10$ days earlier. This is a large gap relative to other SN~I, though similar to SN~2004aw \\citep{Taubenberger2006}; it is caused by metal-line absorption from heavy elements in the outer layers of the ejecta, as is evident from spectra taken after August 1.  A least-squares fit to a $f(t)\\propto(t-t_0)^n$ behavior of our well-sampled early photometry places the explosion date between 10 and 40\\,hr prior to our discovery at the 95\\% confidence level\\footnote{The fireball models can at best give an underestimate on the explosion date, since they do not incorporate the photon diffusion time.}. We cannot rule out the popular quadratic fireball model, but higher values of $n$ are possible as well for larger areas in the fit parameter space (Figure 1, bottom panels).   Spectroscopically, PTF12gzk exhibits high expansion velocities, $\\sim30,000$\\,km\\,s$^{-1}$ (Si~II absorption velocity). Other SNe~Ic with similar velocities are broad-lined SNe~Ic (Figure 4), some of which are associated with GRBs (Woosley \\& Bloom 2006, and references within), while no such association was determined for PTF12gzk (see also SN 2009bb; Soderberg et al. 2009). Most similar is SN 2003lw, a SN associated with a GRB \\citep{Mazzali2006a}. \\begin{figure}[ht] \\centering \\scalebox{0.75}{\\includegraphics{fig7.eps}} \\caption{ {\\scriptsize PTF12gzk has characteristic velocities of a broad-lined SN~Ic. All SNe above the dashed line, besides PTF12gzk, are  GRB-SNe (Xs) or Type Ic-BL with no GRB association (circles), while those below it are normal SNe~Ic (squares). Velocities are obtained through modeling of the spectrum or through direct measurements of the Si~II 6355 \\,\\AA\\ line (SNe 2010bh, 2002ap, 2003jd, 2004aw, and 2007gr). } } \\end{figure} A possible explanation is a burst misaligned with our line of sight, or a failed GRB. Such a scenario is further supported by the host-galaxy characteristics, resembling those of a broad-lined SN~Ic host galaxy. We know of no typical SN~Ic exploding in a host with similar luminosity and oxygen abundance.  The observed relatively narrow lines give a dispersion of $\\Delta v/v \\approx 0.25$, compared to $\\sim1$ in the case of broad-lined SNe~Ic, and may suggest a nonspherical explosion geometry \\citep{Leonard2006}, or that the ejecta mass is high or has a very steep density gradient (a discussion on the effects of a steep density gradient on the LC can be found in Piro \\& Nakar 2012). Late-time, nebular spectra will probe the geometry of the explosion in more detail.  From the Si~II line velocity at peak brightness of PTF12gzk (15,300\\,km\\,s$^{-1}$ from the August 12 spectrum) and SN 2004aw (12,400\\,km\\,s$^{-1}$; Deng et al., in prep.), and the rise time of these two SNe, we use the following scaling relations (Arnett 1982; Mazzali et al. 2009; see also Mazzali et al., in prep.) to estimate the physical properties of PTF12gzk: $\\tau\\approx\\kappa^{1/2}\\mathrm{M}^{3/4}\\mathrm{E}^{-1/4}$ and $v=(2E/M)^{1/2}$, where $\\tau$ is the light-curve rise time, $E$ is the kinetic energy, and $\\kappa$ is the opacity. The derived ejecta mass is 7.5\\,M$_\\odot$ (6--12\\,M$_\\odot$), pointing to a large initial progenitor mass of 25--35\\,M$_\\odot$, though the latter values are highly uncertain \\citep{Mazzali2000}. We derive a kinetic energy of $7.5 \\times 10^{51}$\\,erg [(5--10) $\\times10^{51}$\\, erg]. Using $L_{\\rm max}$ and $t_{\\rm max}$ we get an estimated $^{56}$Ni mass of 0.37\\,M$_\\odot$ (scaling the PTF12gzk light curve to that of SN 2003dh; Mazzali et al. 2003).  Using the $V$-band peak magnitude vs. nickel mass relation presented by Perets et al. (2010), we derive a $^{56}$Ni mass of 0.35\\,M$_\\odot$, in agreement with the results derived from the scaling relations.  These physical properties, as well as the high expansion velocities and the host galaxy, are unlike those of normal SNe~Ic, which typically occur in large hosts and have low ejecta masses, and kinetic energies, (1.7\\,M$_\\odot$, and $10^{51}$\\,erg Drout et al. 2011. For nickel mass in Type Ic SNe see for example Taubenberger et al. 2006, and Sauer et al. 2006). Instead, they are reminiscent of GRB-SNe \\citep{Mazzali2009}.  PTF12gzk is a remarkable example of a SN~Ic in terms of expansion velocities, evolution timescale, the ejected mass, and the kinetic energy released in the explosion. We conclude that these properties point to the explosion of a massive star deficient in H and He, at the higher-mass end of SN~Ic progenitors. This further illustrates the peculiar population of SNe~Ic exploding in dwarf hosts \\citep{Arcavi2010}, as seen also in the case of GRB-SNe and most SLSNe-I.  PTF12gzk demonstrates the advantages of using an untargeted sky survey such as PTF with an extensive network of instruments and telescopes in various wavebands to detect and rapidly characterize unusual cases of cosmic explosions.  \\bigskip \\smallskip  S.B. is supported by a Ramon Fellowship from ISA. A.G. acknowledges support by grants from the ISF, BSF, GIF, Minerva and the EU FP7/ERC. A.V.F. and his group benefit from financial assistance from Gary \\& Cynthia Bengier, the Richard \\& Rhoda Goldman Fund, the Sylvia \\& Jim Katzman Foundation, the Christopher R. Redlich Fund, the TABASGO Foundation, NSF grants AST-0908886 and AST-1211916, and NASA/{\\it HST} grant GO-12530 from STScI (which is operated by the AURA, Inc., under NASA contract NAS 05-26555). P.A.M. and E.P. acknowledge financial support from grants INAF PRIN 2011 and ASI/INAF I/088/06/0. M.I. and Y.J. were supported by the Creative Initiative program of the NRFK. M.M.K. acknowledges Hubble and Carnegie-Princeton Fellowships. D.C.L. is supported by NSF grant AST-1009571. E.O.O. acknowledges the Arye Dissentshik career development chair and a grant from the Israeli MOST.  PTF is a collaboration of Caltech, LCOGT, the Weizmann Institute, LBNL/NERSC, Oxford, Columbia, IPAC, and UC Berkeley. The Liverpool Telescope is operated on the island of La Palma by Liverpool John Moores University in the Spanish Observatorio del Roque de los Muchachos  support from the UK STFC. Construction of the LAIWO camera was supported by the MPIA, GIF, and the ISF. We are grateful for the assistance of the staff at the various observatories used to obtain data. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by JPL, Caltech, under contract with NASA.  \\clearpage"
    },
    "1208/1208.0008_arXiv.txt": {
        "abstract": "In this paper we combine high resolution N-body simulations with a semi analytical model of galaxy formation to study the effects of a possible Warm Dark Matter (WDM) component on the observable properties of galaxies.  We compare three WDM models with a dark matter mass of $0.5$, $0.75$ and $2.0$ keV, with the standard Cold Dark Matter case.  For a fixed set of parameters describing the baryonic physics the WDM models predict less galaxies at low (stellar) masses, as expected due to the suppression of power on small scales, while no substantial difference is found at the high mass end. However these differences in the stellar mass function, vanish when different set of parameters are used to describe the (largely unknown) galaxy formation processes. We show that is possible to break this degeneracy between DM properties and the parameterization of baryonic physics by combining observations on the stellar mass function with the Tully-Fisher relation (the relation between stellar mass and the rotation velocity at large galactic radii as probed by resolved HI rotation curves). WDM models with a {\\it too warm} candidate ($m_{\\nu}<0.75$ keV) cannot simultaneously reproduce the stellar mass function and the Tully-Fisher relation. We conclude that accurate measurements of the galaxy stellar mass function and the link between galaxies and dark matter haloes down to the very low-mass end can give very tight constraints on the nature of DM candidates. ",
        "introduction": "\\label{sec:intro}  The Cold Dark Matter (CDM) paradigm successfully describes the formation of large-scale structure in the Universe (e.g., Springel et al. 2006).  On small scales, the CDM model faces inconsistencies with observations.  For example, CDM model over-predicts the number of satellites in the Milky Way (e.g. Klypin et al. 1999), predicts cuspy density profile for the dark matter halo (e.g.  Moore et al.  1999), and a too large kinematics for massive satellites (Boylan-Kolchin et al.  2011).  These small-scale problems in CDM all arise from the comparison of pure N-body simulations with observational data.  Recently it has been shown that some of those problems could be alleviate by the inclusion of baryonic physics, due to its back-reaction on the properties of dark matter haloes. For example, UV background, reionization and supernova feedback can act together to suppress the formation of dwarf satellites in the Milky Way (e.g. Bullock et al 2001, Macci\\`o et al. 2010), the inner halo density profile becomes more flat when star formation and feedback are included (e.g. Mashchenko et al. 2008, Pontzen \\& Governato 2012, Macci\\`o et al. 2012). However, these baryonic effects are difficult to model and still strongly rely on simple parameterization for complex physical processes (the so called sub-grid physics)   Another possible solution is to suppress the excess of power at small scales in the Dark Matter distribution by changing the properties of the dark matter candidate.  Among many of these different approaches, the Warm Dark Matter (WDM) model is the most intriguing as it naturally preserves the large-scale structure of the CDM model (Alam et al.2002; Zentner \\& Bullock 2002). Many authors have shown that the WDM model can relax some tensions between observations and theoretical predictions using pure N-body simulations (Colin et al. 2000, Bode \\etal 2001, Gotz \\& Sommer-Larsen 2002, Knebe \\etal 2008, Colin \\etal 2008, Tikhonov et al 2009, Lovell et al.2012, Macci\\`o \\etal 2012b).    One possible WDM candidate is a sterile neutrino which exhibit a significant primordial velocity distribution and thus damp primordial inhomogeneities on small scales (e.g. Hansen \\etal 2002, Abazajian \\& Koushiappas 2006, Boyarsky et al. 2009). Limits on the mass of dark matter particles can be obtained from several astrophysical observations: one of the most powerful tool for constraining the matter power spectrum are Lyman-$\\alpha$ forest observations (neutral hydrogen absorption in the spectra of distant quasars, Narayanan \\etal 2000, Viel \\etal 2005, Seljack \\etal 2006). Lyman-$\\alpha$ observations allows the possibility to studying the power spectrum down to small scales and over a large range of redshifts.  Current observations set a lower limit of $m_{\\nu} \\approx 1$ keV (Viel \\etal 2008).  These limits have been confirmed by other observations based on different methods like QSO lensing (Miranda \\& Macci\\`o 2007), luminosity function of high redshift QSOs (Song \\& Lee 2009), dwarf galaxies in the Local Volume (Zavala \\etal 2009), the size of (min)-voids around the Local Group (Tikhovov \\etal 2009) and Milky Way satellites (Macci\\`o \\& Fontanot 2010).   Up to now, most studies based on numerical simulations have focused their attention on the internal structure and kinematics of nearby galaxies in the Warm Dark Matter model, while very few have explored the effect that a warm candidate will have on the properties of the general galaxy population, such as the luminosity function, the stellar to halo mass function and so on.   A notable exception is the recent paper by Menci et al. (2012), although not directly based on N-body simulations in WDM.  In their study the authors combined a Semi Analytical Model for galaxy formation with the Extended Press and Shechter formalism (e.g. Cole \\etal 2000, Lacey \\& Cole 1993) to obtain the properties of galaxy in a WDM model. They found that the WDM model can produce a more flat slope for the faint end of the luminosity function in better agreement with the observations.  They also argue that other galaxy properties better match observations when a warm candidate is used instead than a cold one. While these effects are certainly true, the adoption of a Warm Dark Matter like spectrum has also strong influence on the ratio between dark matter and stellar mass in galaxies. This ratio that can be constrained by several dynamical mass estimators such as the Tully-Fisher relation (Tully \\& Fisher 1977), satellites kinematics, and weak gravitational lensing.  In this paper we combine high resolution N-body simulations for the standard CDM scenario and three WDM candidates $m_{\\nu}=2,0.75,0.5$ keV, with the Semi Analytical Models first presented in Kang et al. (2005), and then successively expanded and improved in Kang \\& van den Bosch (2008) and Kang et al. (2012).  Our aim is to study the impact of different WDM models on galaxy properties and use different observations to try to disentangle the effects of a warm candidate from the different parameterization of the baryonic physics adopted in the SAM.  The final goal is to constrain the properties of the dark matter candidate.  The paper is organized as follows: We introduce the simulations and model in \\S2, show the model predictions in comparison to observations in \\S3, and conclude in \\S4.   \\begin{figure*} \\centerline{\\psfig{figure=fig1.eps,width=0.5\\textwidth}} \\caption{The halo mass  functions from the CDM and WDM simulations. Here the mass functions are slightly different from the usual ones, as we have included the satellite haloes, for which we use the mass at the time of accretions. } \\label{fig:HMFs} \\end{figure*}  \\begin{figure*} \\centerline{\\psfig{figure=fig2.eps,width=0.95\\textwidth}} \\caption{The galaxy stellar mass functions. lines are from our model  with $f_{\\rm c}=0.15$ (see text). The $z=0$ data points are from Li \\& White (2009) and Bell et al.  (2003), and high-z data points are from Marchesini et al. (2009), Fontana et al. (2006).} \\label{fig:SMFs} \\end{figure*} ",
        "conclusions": "\\label{sec:cons}  Recent observational results have challenged the otherwise successful Cold Dark Matter model on small scales. For example the inner density profile of dark matter halo is too concentrated to match the kinematics of satellites around the Milky Way (e.g. Boylan-Kolchin \\etal 2011). Warm Dark Matter models have been suggested as a possible solution since the suppression of power on small scales can in principle alleviate if not solve these issues (e.g. Lovell \\etal 2012).  However, Most of the CDM predictions are based on gravitational only (N-body) simulations, which by construction neglect the effects of baryons and their complicated network and interactions through gas cooling, star formation and feedback that could possibly alter the results of pure DM simulations (e.g. Governato \\etal 2012, Brooks \\etal 2012) It is then important to investigate the prediction of CDM and WDM on the statistical properties of galaxy population, moving beyond a simple DM simulation.   In this paper, we performed high-resolution N-body cosmological simulations for three WDM models with $m_{\\rm \\nu}=0.5$, 0.7 and 2.0 keV respectively, plus a {\\it controlled} CDM model. We couple these simulations with a Semi Analytical Model of galaxy formation to study the impact of a cut-off in the power spectrum on observable quantities such as the stellar mass function.  For a fixed set of parameters describing the baryonic physics, models with low masses for the warm particle ($m_{\\nu}=0.5$ and 0.75 keV), predict less galaxies with stellar mass $M_{\\rm star} <10^{10}M_{\\odot}$ than the current data at $z=0$.  We find that the WDM model with $m_{\\rm \\nu}=2$ keV provides almost identical results as CDM and it is able to successfully reproduce the data. The situation is reversed at higher redshift $z=1.5$, with WDM models in better agreement with the observed stellar mass function than the CDM model (or the WDM model with $m_{\\rm \\nu}= 2$ keV).   However we show these differences in the stellar mass function, are strongly degenerate with the set of parameters used in the Semi Analytical Models to describe the (largely unknown) galaxy formation processes.  By adjusting a single parameter (the cooling rate of low mass haloes) we have been able to get indistinguishable results from all the DM models, both at high and low redshift.  This shows that a single observable (e.g. the stellar mass function) can not constrain the effects of warm component on the galaxy formation process.  In order to break this degeneracy we use independent constraints on the (integrated) star formation efficiency at low masses.  We found that if the WDM models are tuned to  reproduce the present ($z=0$) stellar mass function, the stellar mass for a given halo mass is systematically larger in haloes of masses below ($\\sim 10^{12}M_{\\odot}$). Such scales are not reliably probed by direct methods such as satellite kinematics or weak galaxy-galaxy lensing. In order to probe such low halo masses we use the relation between stellar mass and rotation velocity at large galactic radii, more commonly known as the Tully-Fisher relation. Current data already rule out models with $m_{\\rm \\nu}<0.75$ keV which is in agreement with other limits from large scale structure (e.g. Viel \\etal 2008).  Finally our study shows that by combining measurements of galaxy stellar mass function and stellar mass - halo mass relation down to low galaxy masses ($\\approx 70$km/s) it is possible to obtain very tight constraints on the mass of a possible warm component.  This opens a new window in the search for the nature of the elusive dark matter component of our Universe."
    },
    "1208/1208.0839_arXiv.txt": {
        "abstract": "Complimenting recent work on the effective field theory of cosmological large scale structures, here we present detailed approximate analytical results and further pedagogical understanding of the method. We start from the collisionless Boltzmann equation and integrate out short modes of a dark matter/dark energy dominated universe ($\\Lambda$CDM)  whose matter is comprised of massive particles as used in cosmological simulations. This establishes a long distance effective fluid, valid for length scales larger than the non-linear scale $\\sim$ 10\\,Mpc, and provides the complete description of large scale structure formation. Extracting the time dependence, we derive recursion relations that encode the perturbative solution. This is exact for the matter dominated era and quite accurate in $\\Lambda$CDM also. The effective fluid is characterized by physical parameters, including sound speed and viscosity. These two fluid parameters play a degenerate role with each other and lead to a relative correction from standard perturbation theory of the form $\\sim 10^{-6}c^2\\,k^2/H^2$. Starting from the linear theory, we calculate corrections to cosmological observables, such as the baryon-acoustic-oscillation peak, which we compute semi-analytically at one-loop order. Due to the non-zero fluid parameters, the predictions of the effective field theory agree with observation much more accurately than standard perturbation theory and we explain why. We also discuss corrections from treating dark matter as interacting or wave-like and other issues. \\let\\thefootnote\\relax\\footnotetext{Electronic address: {\\tt mphertz@stanford.edu, mphertz@mit.edu}} ",
        "introduction": "An effective field theory is a description of a  system that captures all the relevant degrees of freedom and describes all the relevant physics at a macroscopic scale of interest. The short distance (``UV\") physics is integrated out and affects the effective field theory only through various couplings in a perturbative expansion in the ratio of  microphysical UV scale/s to the macroscopic (``IR\") scale being probed. This technique has been systematically used in particle physics and condensed matter physics for many years (e.g., see \\cite{Georgi:1994qn,Manohar:1996cq,Kaplan:2005es}), but has not been fully used in astrophysics and cosmology.  The large scale properties of the universe acts as an important application and is in need of careful analysis.  Of current fundamental  importance is to understand the initial conditions, contents, evolution and formation of the universe. It appears to be adequately described by the so-called $\\Lambda$CDM cosmological model in which the matter content of the universe is primarily dark matter and the late time dark energy is adequately described by a cosmological constant. The early universe was dominated by a cosmic plasma in which baryons were tightly coupled to photons leading to so-called baryon-acoustic-oscillations. The evidence for this model comes from a range of sources, including CMB data, lyman-$\\alpha$ forrest, curvature constraints, supernovae type IA, weak lensing, and (of particular importance to the current discussion) structure formation, etc. In this cosmological model, structure formation is primarily driven by the gravitational attraction of dark matter, which led to the clumping of baryons  including stars, galaxies, and clusters of galaxies. This arose from gravitational instabilities of the initial linear density fluctuations that were approximately adiabatic, scale-invariant, and Gaussian (e.g., see \\cite{Peebles:2012mz,Peebles:2009hw,Primack:2006it,CervantesCota:2011pn}).  The power spectrum of large scale structure at late times is corrected from the initial linear input in interesting and important ways. For instance, non-linear effects alter the shape of the baryon-acoustic-oscillations in the power spectrum. The baryon-acoustic-oscillations are a vitally important probe of dark energy as they provide a standard ruler to constrain the cosmological expansion history (e.g., see \\cite{Bassett:2009mm,Weinberg:2012es,Sanchez:2008iw,Wang:2006qt}). In general, one needs a proper understanding of the departures from linear theory in order to constrain fundamental physics, such as dark energy, primordial non-Gaussianity, and other probes of microscopic physics.  There has been a substantial amount of work in the literature to understand  non-linear structure formation in the form of cosmological perturbation theory, including \\cite{Bernardeau:2001qr,Jain:1993jh,Takahashi:2008yk,Shoji:2009gg,Jeong:2006xd,Crocce:2005xy,Crocce:2007dt,Matsubara:2007wj,McDonald:2006hf,Taruya:2007xy,Izumi:2007su,Matarrese:2007aj,Matarrese:2007wc,Nishimichi:2007xt,Peebles:2000yy,Carlson:2009it,Pietroni:2011iz,Tassev:2011ac,Roth:2011ru}. This includes what is usually referred to as ``standard-perturbation-theory\" (SPT). In this approach, the continuity and Euler equations for a pressure-less and non-viscous dark matter fluid (with vanishing stress-tensor) is assumed. These non-linear equations for the dark matter are solved perturbatively around the linear solution and corrections are obtained order by order, usually truncated at the one-loop order. The theory involves integrating $k$-modes in the entire domain $0<k<\\infty$, which involves treating all $k$-modes as perturbative.  The analysis of the current paper stems from the fact that this ``standard\" procedure necessarily has a qualitative and quantitative problem. The density fluctuations are not perturbative beyond a scale $\\knl$, the ``non-linear scale\"; the scale at which density fluctuations are $\\delta\\sim\\mathcal{O}(1)$, where gravitational collapse may occur. This scale is roughly $\\lambda_{NL}\\sim10$\\,Mpc or so. Hence there are two regimes: $k<\\knl$ which is weakly coupled and perturbative, and $k>\\knl$ which is strongly coupled and non-perturbative. For cosmological purposes, we are normally interested in the low $k$-regime. However, these 2 regimes are coupled by non-linearities, so we must be very careful in attempts to describe the low $k$-regime perturbatively. The rigorous and complete method to do this is that of effective field theory. The procedure is to introduce some arbitrary cutoff $\\Lambda$ on the $k$-modes of the fluid. We take this cutoff to be $\\Lambda\\lesssim \\knl$, so that all modes of the fluid are perturbative. This means that the high $k$-modes ($k>\\Lambda$) must be ``integrated out\". In practice this means that these UV modes generate higher order derivative and non-linear corrections to the fluid equations for the low $k$-modes ($k<\\Lambda$). We show that this includes terms such as pressure and viscosity; precisely the terms that are assumed to vanish in SPT. These terms are a real property of the dark matter fluid and they alter the power spectrum in an important and measurable way. These fluid parameters can be determined by matching to the full UV theory, i.e., N-body simulations. This furnishes an effective field theory for dark matter. This is a fluid that only involves weakly coupled modes and is connected to the full microphysical theory through these higher order operators.  Important earlier work on this topic was performed in the very interesting Ref.~\\cite{Baumann:2010tm}, where this basic conceptual foundation was laid out with particular focus on the issue of back-reaction at the scale of the horizon. In our recent work, which we continue here, we (i) focus on sub-horizon scales, (ii) obtain an explicit measurement of fluid parameters, (iii) perform an explicit computation of the power spectrum, and (iv) provide further insight and clarifications. In the present paper we compliment and extend our important recent work in Ref.~\\cite{Carrasco:2012cv} in which the measurement of fluid parameters was performed and the basic framework was put together. Here we develop and describe in detail this effective fluid description of dark matter (and by extension; all matter, since baryons trace dark matter on large scales), which is the complete description of large scale structure formation. In particular, we recapitulate how to extract various fluid parameters from N-body simulations and then solve the effective fluid theory to some desired order in a perturbative expansion which we formulate recursively. We show how to approximately extract the time dependence in a $\\Lambda$CDM universe, which connects in a simple and intuitive way with the previous standard perturbation theory, but highlights the essential differences arising from the fluid parameters. This leads to convenient and quite accurate results. Our basic method and key results are summarized in the following discussion:  If we assume that the matter content of the universe is dominated by non-relativistic matter, primarily dark matter evolving under Newtonian gravity, we can smooth the corresponding collisionless Boltzmann equation in an expanding FRW background. This generates the usual continuity and Euler equations. An important point stressed in Refs.~\\cite{Baumann:2010tm,Carrasco:2012cv} is that the latter includes an effective stress-tensor $[\\tau^{ij}]_\\Lambda$ that is sourced by the short-modes $\\delta_s$. By defining the effective stress-tensor by its correlation functions, it can be expanded in terms of the long density  $\\delta_l$ and velocity $v_l^i$ fields as \\bea [\\tau^{ij}]_\\Lambda\\amp=\\amp \\delta^{ij}p_b+\\rho_b\\Bigg{[} c_s^2\\,\\delta^{ij}\\delta_l-{c_{bv}^2\\over Ha}\\delta^{ij}\\,\\partial_k v_l^k\\nonumber\\\\\\amp-\\amp{3\\over4}{c_{sv}^2\\over Ha}\\left(\\partial_jv_l^i+\\partial_iv_l^j-{2\\over3}\\delta^{ij}\\,\\partial_kv_l^k\\right) \\Bigg{]}+\\ldots \\eea The individual parameters $c_s^2$ and $c_v^2\\equiv c_{sv}^2+c_{bv}^2$ are degenerate with each other at the one-loop order since we only track the growing modes, and degenerate with other parameters at higher loop order. While the shear viscosity parameter $c_{sv}^2$ affects the vorticity, which is a somewhat small effect. As analyzed in Ref.~\\cite{Carrasco:2012cv} one can directly evaluate the stress-tensor from the microphysical theory, i.e., from N-body simulations to extract such parameters. For  smoothing scale $\\Lambda=1/3$\\,[h/Mpc] at $z=0$ it is found $c_s^2+f c_v^2\\approx 9\\times 10^{-7}c^2$, ($f$ is the logarithmic derivative of the growth function). This direct measurement can also be obtained from matching to the power spectrum at some renormalization scale, resulting in a consistent value and a positive check on the validity of the theory.  We establish recursion relations for the density fluctuations and velocity field, allowing us to insert this measured value of the fluid parameter and obtain correlation functions. These parameters carry $\\Lambda$ dependence which balances the $\\Lambda$ dependence of the cutoff on the loops. If we take $\\Lambda$ to large values the fluid parameters approach a finite quantity, representing the finite error made in the standard perturbation theory. In particular the fluid parameters provide the following relative correction to the power spectrum (suppressing the time dependence and $\\mathcal{O}(1)$ factors) \\beq {\\delta P(k)\\over P_{L}(k)}\\sim -{10^{-6}\\,c^2\\,k^2\\over H^2} \\label{Pcorr}\\eeq where $P_{L}(k)$ is the linear power spectrum. % Note that the pressure and viscosity act together to reduce the power spectrum by acting oppositely to gravity. This simple, but entirely real and rigorous correction to the power spectrum is essential to explain the observed shape of the baryon-acoustic-oscillations in the power spectrum relative to standard theory.  The outline of the paper is the following: In Section \\ref{EffectiveFluid} we describe the basic theoretical setup. Operating in the Newtonian approximation in an expanding universe, we smooth the Boltzmann equation to obtain an effective fluid for cold dark matter. We recapitulate how to match its parameters to the microphysical results from N-body simulations. In Section \\ref{PerturbationTheory} we solve the theory recursively for a matter dominated universe by extracting the time dependence, and lift this to $\\Lambda$CDM in an approximate way also. This allows us to semi-analytically derive the power spectrum at one-loop order. In Section \\ref{PowerNumResults} we present our numerical results for the power spectrum and compare to linear theory and standard perturbation theory. In Section \\ref{Discussion} we discuss the fluid's parameters, corrections from collisions, wave-like behavior, higher order moments, and the velocity field. Finally, in Section \\ref{Summary} we summarize the effective field theory and its role in  cosmology. ",
        "conclusions": "\\label{Discussion}  In this section we briefly mention some interesting issues surrounding the effective fluid, including its Reynolds number, degeneracy of parameters, the inclusion of collisions or wave-like behavior, higher order moments, and the velocity field.  \\subsection{Reynolds number} For viscous fluids there is a famous dimensionless number which captures its tendency for laminar or turbulent flow; the ``Reynolds number\". The Reynolds number is defined as \\beq R_e\\equiv {\\rho \\, v L\\over\\eta} \\eeq where $\\eta$ is shear viscosity, $\\rho$ is density, $v$ is a characteristic velocity, and $L$ is a characteristic length scale. This is \\beq R_e\\sim {H v L\\over c_{sv}^2}\\sim {H^2a^2\\over c_{sv}^2 k^2}\\delta\\lesssim 10 \\eeq Hence the Reynolds number is not very large, and the system is therefore not turbulent. Furthermore, if we were to estimate the viscosity by Hubble friction, then we would have $R_e\\sim\\delta$ and so the Reynolds number would be even smaller in the linear or weakly non-linear regime.  For cosmological parameters $\\rho_b\\sim 3\\times 10^{-30}$\\,[g/cm$^3$], $H=70$\\,[km/s/Mpc], and if we take a plausible value for the shear viscosity of $c_{sv}^2\\sim 2\\times 10^{-7}c^2$, then the viscosity coefficient is found to be \\beq \\eta\\sim 20\\, \\mbox{Pa\\,s} \\eeq which is perhaps surprisingly not too far from unity in SI units. (For instance, it is somewhat similar to the viscosity of some everyday items, such as  chocolate syrup.) A proper measurement of $c_{sv}^2$ would come from a detailed measurement of vorticity; a point that we will return to in the following subsection.  \\subsection{Degeneracy in Parameters}\\label{Degeneracy}  Earlier we discussed the individual parameters: the sound speed $c_s^2$, shear viscosity $c_{sv}^2$, and bulk viscosity $c_{bv}^2$. We showed that by taking the curl of the Euler equation we obtained an equation for vorticity that involves $c_{sv}^2$. Hence a careful analysis of vorticity could reveal the value of $c_{sv}^2$ - such a value would enter our discussion of ``Reynolds number\" of the previous subsection, although not the bulk of this paper. Since the vorticity is rather small, this would be non-trivial to measure, though possible.  On the other hand, by taking the divergence of the Euler equation we obtained coupled equations for $\\theta_l$ and $\\delta_l$ that involves the sound speed $c_s^2$ and the combination of viscosity $c_v^2=c_{sv}^2+c_{bv}^2$. One could try to measure these two parameters independently using eqs.~(\\ref{csx},\\,\\ref{cvisx}) from N-body simulations. However one should be careful as to how to interpret this result. At the one-loop level, we saw that it was only a certain linear combination that appeared in the result, namely $c_s^2+f\\,c_v^2$. In other words, the two parameters appear in a degenerate way at one-loop. This degeneracy would be broken at higher loop order. However, to be self-consistent one should then also include new couplings (for instance, representing higher derivative operators in the stress-tensor expansion) which would also enter, leading to a new constraint and a new type of degeneracy with the new parameters.  This reason for this degeneracy is the following: In a universe in which one could  track the full evolution of the initial state, one would observe that $c_s^2$ and $c_v^2$ affect the one-loop theory differently. However, in the real universe, there is a growing mode and a decaying mode. In practice, one does not track the decaying mode, only the growing mode, as studied here in this paper. For this single mode the parameters enter in a special linear combination.   \\subsection{Interactions}\\label{Collisions} In this paper, we have treated dark matter as being comprised of collisionless particles, interacting only through gravity. Of course we expect that there are also some finite non-gravitational interactions. These effects can be treated perturbatively in the effective field theory by identifying the relevant length scale, which is the mean free path between scattering.  As an example, for WIMPs the scattering cross section is roughly $\\sigma\\sim g^2/m_W^2$, where $g\\sim 0.1$ is a dimensionless coupling and $m_W\\sim 100$\\,GeV is of order the weak scale. The mean free path between scatterings is \\beq \\lmfp={1\\over n\\,\\sigma}\\sim {m_W^3\\over\\rho_b\\, g^2} \\eeq The background density, in Planck units, is $\\rho_b\\sim 10^{-120}\\mpl^4$. Hence, \\beq \\lmfp\\sim 10^{17}d_H\\left(0.1\\over g\\right)^2\\left(m_W\\over100\\,\\mbox{GeV}\\right)^3 \\eeq where $d_H=1/H_0$ is the present Hubble length. As another example, for QCD-axions scattering due the $~\\phi^4$ term in the potential $V(\\phi)=\\Lambda_{bcd}^4(1-\\cos(\\phi/F_{PQ}))$, the mean free path is much larger still.  On the other hand, gravity introduces a non-linear scale of the order $\\lnl\\sim 10^{-4}\\,d_H$ (which can be thought of as the mean free path between gravitational scattering of a point particle off a large non-linear clump). Since the collisional mean free path in these examples satisfies $\\lmfp\\gg\\lnl$ it can be ignored at first approximation. Though in principle it can be included perturbatively in the effective field theory, but suppressed by a hierarchy $\\lnl/\\lmfp$, which may be of interest for some highly non-standard dark matter candidates.  \\subsection{Wave-like Behavior}\\label{Wave}  In this paper we have treated the dark matter as comprised of classical point-like particles. This obviously ignores its quantum mechanical wave-like behavior. For most dark matter candidates, such as a typical WIMP with a weak scale mass, the de Broglie wavelength is extremely small and ignorable. For extremely light (pseudo)-scalars, such as axions, it is conceivable that their de Broglie wavelength is large and relevant.  For a classical scalar field and also for a Bose-Einstein condensate \\cite{Sikivie:2009qn}, one can show that in the linear theory, there is a correction to the pressure of the form \\beq \\delta p=-{\\hbar^2\\over 4\\,a^2 m^2}\\nabla^2\\rho \\eeq this provide a contribution to a type of scale dependent sound speed $\\delta c_s^2\\sim \\hbar^2 k^2/(a^2m^2)$. Now recall that the characteristic correction from the sound speed is $\\sim c_s^2k^2/(H^2a^2)$. This means that a rough estimate for the dimensionless correction from the quantum character of the particles is (at $z=0$) \\beq \\mbox{quantum correction} \\sim {\\hbar^2 k^4\\over m^2H_0^2} \\eeq The relative size of this contribution obviously depends on the mass of the particle $m$. It is important to note that in the point-particle treatment, the mass $m$ dropped out of all results. But by including such UV physics, we gain more sensitivity in the effective field theory to such physical parameters.  The dark matter particle mass $m$ can in principle be very small. For instance, in the so-called string axiverse it is suggested that there may be a range of extremely light axions \\cite{Arvanitaki:2009fg}; one of which could provide the bulk of the dark matter (though there are important constraints from isocurvature bounds on light axions \\cite{Hertzberg:2008wr,Hamann:2009yf}, while the classic axion window is still very promising \\cite{Hertzberg:2012zc,Erken:2011dz}). Here we would like to mention that the mass of a dark matter particle presumably cannot be arbitrarily small because its de Broglie wavelength $\\lambda_{dB}\\sim \\hbar/(m v)$ would then smear it out over scales larger than that of a galaxy $L_{gal}$, and yet we know dark matter clumps on galactic scales. By imposing $\\lambda_{dB}<L_{gal}$ this gives the bound \\beq m>{\\hbar\\over L_{gal}v} \\eeq Hence we have a bound on the dimensionless correction from the wave-like character of light scalars as \\beq \\mbox{quantum correction}<{k^4L_{gal}^2v^2\\over H_0^2} \\eeq where $v$ is a characteristic dispersion velocity associated with the dark matter. By estimating $v\\sim 10^{-3}\\,c$, then we can estimate $H_0/v\\sim k_{NL}$, leading to the rough bound \\beq \\mbox{quantum correction}<{k^4L_{gal}^2\\over k_{NL}^2} \\eeq Since $L_{gal}$ is much smaller than the non-linear scale (for instance, $L_{gal}$ may be as small as dwarf galaxy size) we see that the quantum correction must always be very small in the regime in which the effective field theory is valid (i.e., $k<k_{NL}$).  \\subsection{Higher Order Moments}  In principle, one can study higher order moment of the Boltzmann equation. In Section \\ref{EffectiveFluid} we considered the zeroth moment (continuity) and first moment (Euler), and then built a derivative expansion for the effective stress-tensor that appears on the right hand side of the Euler equation. The stress-tensor involves two contributions: kinetic and gravitational. The kinetic piece $\\kappa_l^{ij}$ is includes the second moment of the velocity distribution (minus the long modes), and so it evolves under the second moment of the Boltzmann equation. The trace of the kinetic part of the stress-tensor is proportional to a type of ``kinetic temperature\" $T\\sim\\kappa_l/\\rho_l$. Although the system is not in thermal equilibrium, so this name is only by analogy to classical systems which are.  However, since the stress-tensor also includes the gravitational piece $w^{ij}$, we do not have an evolution equation for the full stress-tensor. Parametrically these two contributions are of the same order (for instance they cancel each other in the virial limit). So this requires the use of a derivative expansion to capture the effects of these higher order moments and interactions, in the effective field theory sense.  \\subsection{Velocity Field}  Let us also make some comments on the computation of correlation functions, such as the two-point, involving the smoothed velocity field. In Section \\ref{CutoffIndependence} we demonstrated how the cutoff dependence in the expansion for $\\delta_l$ cancels out when we form the two-point correlation function for density, which involved the fluid parameter's canceling the $\\Lambda$ dependence of the $P_{13}$ loop. However, once the fluid counter-terms are introduced to cancel this dependence, then they cannot also be used to cancel the cutoff dependence appearing in the loops of correlation functions of other fields. In particular, consider the velocity field. Recall that it is defined by the ratio $v^i_l=\\pi_l^i/\\rho_l$. If we Fourier transform this, then examine the regime $k<\\Lambda$, we are still left with $\\Lambda$-dependence, even though both $\\rho_l$ and $\\pi^i_l$ are cutoff independent for $k<\\Lambda$. This is similar to certain kinds of non-linear objects that one might define to describe pions, such as a bi-linear in the quark fields, which depends explicitly on the cutoff. This means that the velocity field $v_l^i$ is inherently cutoff dependent even in the $k\\ll\\Lambda$ regime, while $\\delta_l$ is not."
    },
    "1208/1208.3467_arXiv.txt": {
        "abstract": "The radio--loud quasar SDSS J102623.61+254259.5, at a redshift $z$=5.3, is one of the most distant radio--loud objects. Since its radio flux exceeds 100 mJy at a few GHz, it is also one of the most powerful radio--loud sources. We  propose that this source is a blazar, i.e. we are seeing its jet at a small viewing angle. This claim is based on the spectral energy distribution of this source, and especially on its strong and hard X--ray spectrum, as seen by {\\it Swift}, very typical of powerful blazars. Observations by the Gamma--Ray Burst Optical/Near--Infrared Detector (GROND) and by the Wide--field Infrared Survey Explorer (WISE)  allow to establish the thermal nature of the emission in the near IR--optical band. Assuming that this is produced by a standard accretion disk, we derive that it emits a luminosity of $L_{\\rm d}\\simeq9\\times10^{46}$ erg s$^{-1}$ and that the black hole has a mass  between 2 and 5 billion solar masses. This poses interesting constraints on the mass function of heavy ($>10^9 M_\\odot$) black holes at high redshifts. ",
        "introduction": "\\label{sec-intro}  Blazars are Active Galactic Nuclei (AGN) whose relativistic jets are seen at a small angle from our line of sight, smaller than $1/\\Gamma$, where $\\Gamma$ is the bulk Lorentz factor of the emitting plasma. The radiation emitted by their jets is therefore strongly boosted, making them visible at high redshifts. The most distant blazar known is Q0906+6930 (Romani et al. 2004; Romani 2006) with a redshift $z=5.47$. The spectral energy distribution (SED) of this source reveals both the thermal (i.e. strong optical emission lines and continuum) and boosted non--thermal (dominating in all bands but the near IR and optical) components. This is a common characteristic of powerful blazars at high redshift: the non thermal emission is characterized by two broad humps, peaking at lower frequencies as the bolometric luminosity increases, making the underlying thermal emission visible in powerful blazars (since it stands between the two non--thermal humps; Ghisellini et al. 2010a; 2010b). The high energy hump, in these sources, peaks below 100 MeV, disfavoring its detection by the Large Area Telescope (LAT) instrument onboard the {\\it Fermi} satellite (Atwood et al. 2009). Instruments observing in the hard X--ray band can instead have more chances to detect them, and in fact the Burst Alert Telescope (BAT) onboard the {\\it Swift} satellite (Gehrels et al. 2004) have detected blazars up to larger redshifts than LAT (compare Ajello et al. 2009 and Ackermann et al. 2011). Their X--ray spectrum is hard [i.e. $\\alpha_X \\lsim 0.5$, assuming $F(\\nu) \\propto \\nu^{-\\alpha_{x}}$], and this, together with a relatively strong X--ray to optical flux ratio, can be taken as a signature of the blazar nature of the source.  In this letter we suggest that SDSS J102623.61+254259.5, a radio--loud AGN at $z=5.3$, is a blazar, i.e. the viewing angle is smaller than $1/\\Gamma$. If so, then this source indicates the presence, at approximately the same redshift, of  $\\sim 2\\Gamma^2$  other sources pointing in other directions, having similar intrinsic physical properties, including the black hole mass. Note that the flatness of the radio spectrum is not sufficient to guarantee the blazar nature of SDSS J102623.61+254259.5, since slightly misaligned sources (with respect to $1/\\Gamma$) still show a flat radio spectrum, especially at relatively large radio frequencies (at $z=5.3$, the observed 1.4 GHz flux corresponds to a rest frame frequency of $\\sim$9 GHz), where the radio lobe steep radio emission does not contribute. Evidences for its blazar nature are its very large radio--loudness and very large radio luminosity. Observed by {\\it Swift}, it revealed a strong and hard X--ray flux, confirming its blazar nature. Observations performed by the Gamma--Ray Burst Optical/Near--Infrared Detector (GROND; Greiner et al. 2008) and by the WISE satellite (Wright et al. 2010), together with the Sloan Digital Sky Survey (SDSS; York et al. 2000) spectrum, allowed to constrain the properties of the thermal emission.   In this work, we adopt a flat cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{\\rm M}=0.3$.     \\begin{table*} \\centering \\begin{tabular}{llllllll} \\hline \\hline ~  &$g^\\prime$ &$r^\\prime$ &$i^\\prime$ &$z^\\prime$ &$J$ &$H$ &$K_{s}$ \\\\ \\hline $\\lambda_{\\rm eff}$ (\\AA)   &4587 &6220 &7641 &8999 &12399 &16468 &21706 \\\\ AB magnitude      &$<$24.0 &22.07$\\pm$0.08 &19.97$\\pm$0.03 &19.86$\\pm$0.04 &19.50$\\pm$0.10 &19.23$\\pm$0.15 &19.07$\\pm$0.21 \\\\ \\hline \\hline \\end{tabular} \\vskip 0.4 true cm \\caption{GROND AB observed magnitudes of B2 1023+25, taken 2012 April 16th (magnitudes not corrected for Galactic extinction). The first raw gives the effective wavelength of the filter. } \\label{grond} \\end{table*} ",
        "conclusions": "\\label{sec-discussion}   In this letter we propose that the radio--loud, high redshift quasar B2 1023+25 is a blazar. It is the second most distant blazar known up to now, with redshift $z=5.3$, that corresponds to an age of the Universe of $10^9$yrs.  Both its thermal and non--thermal components are very luminous. In agreement with the blazar sequence (Fossati et al. 1998) the two broad non--thermal humps peak at small frequency. In particular, the hard X--ray spectrum and the upper limit in the $\\gamma$--ray band constrain the high energy component to peak in the MeV region of the spectrum. Therefore this source, along with the similar other powerful blazars, should have a relatively large hard X--ray flux and would have been an ideal target for planned hard X--ray instruments, such as {\\it EXIST}, {\\it Symbol--X} and {\\it NHXM}, and it is a good target to be observed with the only orbiting satellite with focussing hard X--ray telescopes, i.e. {\\it NuStar} (Harrison et al., 2010). The great sensitivity over its energy range [5--80 keV] enables it to detect the hard X--ray spectrum of this source, even if it cannot directly observe the peak of the high energy hump.  \\vspace{0.3 cm}  Thanks to the good data coverage in the IR--optical range, given by the combination of the 7 simultaneous optical--near IR filters of GROND coupled with WISE and SDSS data, we were able to estimate the central black hole mass, that turned out to be between 2 and 5 billion solar masses. This relatively small range compares well with the virial estimates in general (that have at least a factor $\\sim$4 of uncertainties, see e.g. Vestergaard \\& Peterson 2006), and, in particular, it is the only method applicable for this source, where only a partially absorbed Ly$\\alpha$ line is visible (and virial methods are not yet well calibrated in this case). Near--IR spectroscopy, and hence the detection of the broad CIV and MgII lines, would help to improve the virial estimate of the black hole mass and be a valuable cross--check of our estimate, based on the disk--continuum.   It is remarkable to find radio--loud AGN hosting BH that massive at an age of the Universe of only $10^9$yrs. The discovery that B2 1023+25 is a blazar, together with that one of Q0906+693, in fact, implies the existence of an entire family of radio--loud, very massive quasars in the early Universe, whose jet is not aligned with our line--of--sight. Moreover, it is possible to roughly estimate the number of these objects and their spatial density.  The comoving density of heavy black holes at high redshifts of radio--loud sources has been studied by Volonteri et al. (2011), based on the 3 years BAT catalog and the blazar luminosity function, in hard X--rays, derived by Ajello et al. (2009), and modified (beyond $z=4.3$) by Ghisellini et al. (2010a). In the latter paper the observational constrain on the blazar density with black holes heavier than $10^9 M_\\odot$ in the redshift bin $5<z<6$ was based on the detection of only one object: Q0906+6930. Assuming it was the only blazar in the entire sky in this redshift bin, Ghisellini et al. (2010a) derived a comoving density of $2.63\\times 10^{-3}$ Gpc$^{-3}$ of blazars hosting an heavy black hole in the $5<z<6$ bin (see Fig. 15 in that paper).  B2 1023+25 was selected in the SDSS catalog covered by FIRST observations, and the common area of the sky of these two surveys is 8770 square degrees. Its existence in this portion of the sky implies that the space density of blazars hosting a $M_{\\rm BH}>10^9 M_\\odot$ black hole is a factor $40000/8770$ greater than estimated previously (increasing the corresponding comoving density to $1.2\\times 10^{-2}$ Gpc$^{-3}$). In turn, the density of all radio--loud sources hosting a $M_{\\rm BH}>10^9 M_\\odot$ black hole must be a factor $2\\Gamma^2$ larger, namely $\\sim 4.7(\\Gamma/14)^2$ Gpc$^{-3}$ in the $5<z<6$ bin. Note that this is a lower limit, since other high redshift blazars could exist, but not yet identified as such.  Moreover, assuming that the ratio of radio--loud to radio--quiet AGN is $\\sim 0.1$, the SMBHs with $M_{\\rm BH}>10^9M_\\odot$ have to be {\\it at least} $10^4$ just after $10^9$ yrs from the birth of the Universe."
    },
    "1208/1208.3698_arXiv.txt": {
        "abstract": "We investigate how radiative feedback from the first stars affects the assembly of the first dwarf galaxies. To this end we perform cosmological zoomed smoothed particle hydrodynamics simulations of a dwarf galaxy assembling inside a halo reaching a virial mass $\\sim 10^9 \\Msun$ at $z = 10$. The simulations follow the non-equilibrium chemistry and cooling of primordial gas and the subsequent conversion of the cool dense gas into massive metal-free stars. To quantify the radiative feedback, we compare a simulation in which stars emit both molecular hydrogen dissociating and hydrogen/helium ionizing radiation with a simulation in which stars emit only molecular hydrogen dissociating radiation, and further with a simulation in which stars remain dark. Photodissociation and photoionization exert a strong negative feedback on the assembly of the galaxy inside the main minihalo progenitor. Gas condensation is strongly impeded, and star formation is strongly suppressed in comparison with the simulation in which stars remain dark. The feedback on the gas from either dissociating or ionizing radiation implies a suppression of the central dark matter densities in the minihalo progenitor by factors of up to a few, which is a significant deviation from the singular isothermal density profile characterizing the dark matter distribution inside the virial radius in the absence of radiative feedback. The evolution of gas densities, star formation rates, and the distribution of dark matter becomes insensitive to the inclusion of dissociating radiation in the late stages of the minihalo assembly, and it becomes insensitive to the inclusion of ionizing radiation once the minihalo turns into an atomically cooling galaxy. The formation of a rotationally supported extended disk inside the dwarf galaxy is a robust outcome of our simulations not affected by the inclusion of radiation. Low-mass galaxies in the neighborhood of the dwarf galaxy show a large scatter in the baryon fraction which is driven by radiative feedback from sources both internal and external to these galaxies. Our estimates of the observability of the first galaxies show that dwarf galaxies such as simulated here will be among the faintest galaxies the upcoming {\\it James Webb Space Telescope} will detect. Our conclusions regarding the structure and observability of the first galaxies are subject to our neglect of feedback from supernovae and chemical enrichment as well as to statistical uncertainties implied by the limited number of galaxies in our simulations. ",
        "introduction": "The birth of star-forming galaxies a few hundred million years after the Big Bang marks an important milestone in the history of our universe. The spectacular images returned by the {\\it Hubble Space Telescope} ({\\it HST}) in the past few years have already allowed us to probe into the first billion year of the universe. The last few years have also seen the development of an observationally testable theory of the formation of the first galaxies. As ongoing and upcoming observations, such as with the {\\it James Webb Space Telescope} ({\\it JWST}), are about to push to ever earlier times, approaching the epoch of the very first galaxies, this theory is put to ever more stringent tests (for a review see, e.g., \\citealp{D:2012}). \\par Both analytical arguments (e.g.,\\citealp{Tegmark:1997}; \\citealp{Naoz:2006}; \\citealp{Tseliakhovich:2011}) and simulations (e.g., \\citealp{Abel:2002}; \\citealp{Bromm:2002}; \\citealp{Stacy:2011}) suggest that the first stars have formed at $z \\gtrsim 30$ inside dark matter minihalos with virial temperatures $\\gtrsim 10^3\\K$, corresponding to halo masses $\\sim 10^5-10^6 \\Msun$ (for a review, see \\citealp{Bromm:2004}). The metal-free primordial gas inside these minihalos cools and condenses to reach the high densities needed to form stars primarily through the radiative de-excitation of rovibrationally excited molecular hydrogen (e.g., \\citealp{Haiman:1996b}; \\citealp{Abel:1997}). As the minihalos grow in mass, their virial temperatures increase and, after reaching $\\sim 10^4 \\K$, become sufficiently large to collisionally excite atomic hydrogen. This gives birth to the first atomically cooling galaxies with typical masses $\\sim 10^7 -10^8 \\Msun$ at $z \\gtrsim 15$ (e.g., \\citealp{Oh:2002}; \\citealp{Wise:2007b}; \\citealp{Greif:2008}; for a review see \\citealp{Bromm:2011}). The first atomically cooling galaxies then evolve into the first dwarf galaxies with characteristic masses $\\gtrsim 10^9 \\Msun$ (e.g., \\citealp{Mashchenko:2008}; \\citealp{Wise:2012}). The aim of the current work is to investigate the assembly of such galaxies under the radiative feedback from the first stars. The mass-scale marked by the first dwarf galaxies is closely related to a number of key open issues, some of which are outlined below. \\par \\par {\\it Suppression of star formation by radiative feedback. } The radiation emitted by the first stars has a profound effect on subsequently forming stars and galaxies (for a comprehensive overview, \\citealp{Ciardi:2005}). Radiation in the Lyman-Werner (LW) bands dissociates molecular hydrogen, the main coolant in minihalos (e.g., \\citealp{Haiman:1997}). Hydrogen-ionizing radiation heats the gas inside the first halos and the intergalactic medium (IGM). The associated increase in pressure drives the gas outside halos with virial temperatures $\\lesssim 10^4 \\K$, suppressing star formation in both minihalos and the first atomic cooling halos (e.g., \\citealp{Thoul:1996}; \\citealp{Barkana:1999}). The increased pressure in the IGM impedes the accretion of gas onto these low-mass halos, an effect known as Jeans-filtering (e.g., \\citealp{Shapiro:1994}; \\citealp{Gnedin:1998}; \\citealp{Okamoto:2008}). In contrast, star formation inside the first dwarf galaxies should be more robust to this negative radiative feedback as their deeper gravitational potentials allow them to hold on to their gas more strongly. \\par {\\it The sources of reionization.} The first galaxies are thought to have started the reionization of the universe, which is the transformation of the cosmic hydrogen from its early neutral to its present ionized state that occurred during the first billion year after the Big Bang (for reviews see, e.g., \\citealp{Barkana:2001}; \\citealp{Furlanetto:2006}; \\citealp{Meiksin:2009}). However, whether galaxies could sustain reionization and drive it to completion is a question of significant debate (e.g., \\citealp{Bouwens:2012}; \\citealp{Finkelstein:2012}). The largest uncertainties are related to our poor knowledge of the escape fraction, i.e., the fraction of ionizing photons that leave the galaxies unabsorbed and are thus available to reionize the IGM, and of the abundance and ionizing luminosities of low-mass galaxies too faint to be detected in current surveys (e.g., \\citealp{Robertson:2010}). The most recent determinations of the UV luminosity density at $z \\gtrsim 7$ suggest that a significant contribution from faint, yet to be observed, low-mass galaxies is likely needed to sustain reionization in the recombining gas (\\citealp{Finkelstein:2012}; see \\citealp{KuhlenFG:2012} for a comprehensive discussion). Because of their increased robustness against stellar feedback, dwarf galaxies are among the low-mass galaxies expected to be especially efficient sources of reionization (e.g.,\\citealp{Choudhury:2007}; \\citealp{Raicevic:2011}). \\par {\\it The origin of the Milky Way (MW) satellites.}  Dissipationless simulations of dark matter subhalos around MW-like galaxies imply the existence of a large population of satellite galaxies only few of which are currently observed. A number of solutions have been offered to explain this missing satellite problem (\\citealp{Moore:1999}; \\citealp{Klypin:1999}; for reviews see, e.g., \\citealp{Bullock:2010}; \\citealp{Ricotti:2010}; \\citealp{Mayer:2010}), including the suppression of star formation in low-mass halos by stellar feedback and reionization (e.g., \\citealp{Ricotti:2005}; \\citealp{Salvadori:2009}; \\citealp{Bovill:2009}; \\citealp{Sawala:2012}; \\citealp{Brooks:2012}), the transformation of low-mass halos by tidal forces and ram-pressure stripping upon their entry in the MW virial region (e.g., \\citealp{Mayer:2007}), modifications of the cold dark matter structure formation paradigm (e.g., \\citealp{Lovell:2012}), and others (e.g., \\citealp{Bovy:2012}; \\citealp{Vera-Ciro:2012}; \\citealp{Wang:2012}). However, current theories still struggle to explain the abundance and properties of the observed MW satellites across the luminosity range, from the recently discovered ultra-faint to the long-known classical satellites (e.g, \\citealp{Strigari:2008}; \\citealp{Bovill:2011}; \\citealp{Boylan:2011}). The most massive of the simulated MW dark matter subhalos have progenitors with masses $10^8-10^{10} \\Msun$ at $z \\gtrsim 6$ (e.g., \\citealp{Boylan:2012}). This suggests an intimate relation with the first dwarf galaxies, and renders investigations into these objects a promising tool to understand the origin of structure in galaxies such as the MW. \\par {\\it The first disk galaxies.} Simulations of the first atomically cooling galaxies, i.e., galaxies inside halos with masses $10^7-10^8\\Msun$ at redshifts $z \\sim 15$, reveal a highly irregular morphology of the halo gas (e.g., \\citealp{Wise:2007b}; \\citealp{Greif:2008}; \\citealp{Regan:2009}). On the other hand, simulations of galaxies inside halos with larger masses $\\gtrsim 10^9 \\Msun$ and at lower redshifts $z \\lesssim 6-10$ often find the halo gas organized in rotationally supported disks (e.g., \\citealp{Mashchenko:2008}; \\citealp{Pawlik:2011a}; \\citealp{Romano:2011}; \\citealp{Wise:2012}). These findings suggest a dwarf-size mass scale $\\sim 10^8-10^{10} \\Msun $ for the transition to disk-like morphologies, and the emergence of the first disk galaxies at $z \\gtrsim 6$. Physical processes to imprint such a scale include the turbulence generated by the cold inflow of gas along filaments which characterizes gas accretion by the first atomic cooling halos (e.g., \\citealp{Wise:2007b}; \\citealp{Wise:2008c}; \\citealp{Greif:2008}), and stellar feedback (e.g., \\citealp{Kaufmann:2007}). Whether the first halos may host disks is an important open issue, affecting estimates of, e.g., the escape of ionizing photons into the IGM (e.g., \\citealp{Gnedin:2008}; \\citealp{Conroy:2012}), or the ability of massive black holes to accrete gas and grow (e.g., \\citealp{Eisenstein:1995}; \\citealp{Koushiappas:2004}; \\citealp{Lodato:2006}; \\citealp{Petri:2012}). \\par {\\it The faintest galaxies JWST will see.} The faintest $z \\gtrsim 6$ galaxies {\\it HST} has so far revealed have estimated stellar masses $\\gtrsim 10^8 \\Msun$ (\\citealp{Labbe:2010}; \\citealp{Finkelstein:2010}; \\citealp{Curtis:2012}). Future observations with upcoming telescopes such as the {\\it JWST} will allow to search for galaxies down to still lower stellar masses and out to higher redshifts, thus promising to test our theories of the formation of the first stars and galaxies. However, most studies agree that even {\\it JWST} will not be sensitive enough to detect the stellar radiation emitted from inside the minihalos and the first atomic cooling halos (e.g., \\citealp{Haiman:1998}; \\citealp{Oh:2001}; \\citealp{Ricotti:2008}; \\citealp{Johnson:2009}; \\citealp{Zackrisson:2012}; \\citealp{Rydberg:2012}). {\\it JWST} may detect the stellar radiation from some of these objects if they are gravitational lensed (e.g., \\citealp{Zackrisson:2012}). But most of the stellar light collected by {\\it JWST} from high redshifts is expected to come from dwarf galaxies more massive than the first atomically cooling galaxies (e.g., \\citealp{Johnson:2009}; \\citealp{Pawlik:2011a}). \\par Motivated primarily by the exciting prospects for observations with the upcoming {\\it JWST}, we have previously presented cosmological simulations of a dwarf galaxy assembling inside a halo reaching $10^9\\Msun$ at $z = 10$ (\\citealp{Pawlik:2011a}). The simulations were performed using the Smoothed Particle Hydrodynamics (SPH) technique and achieved high resolution by zooming in a select region around the galaxy. Following the non-equilibrium chemistry and cooling of primordial gas, the simulations tracked the evolution of the dwarf galaxy starting from before its birth inside a minihalo. An intriguing outcome was the formation of a rotationally supported extended disk just prior to the final simulation redshift. However, our previous simulations did not account for star formation and the associated feedback. As explained above, stellar feedback has the potential to significantly affect the assembly of the gas inside low-mass halos. \\par In this study we present a new set of simulations similar to our previous simulations, but extending them by including star formation and radiation. We focus on the radiative feedback from LW and ionizing radiation.\\footnote{We will use the terms LW radiation and dissociating radiation interchangeably.} To judge the impact of radiative feedback we will compare a simulation that includes both LW and ionizing radiation with a simulation that includes only LW radiation and further with a simulation in which no radiation is emitted. Note that the simulations do not account for supernova (SN) feedback or metal enrichment, a limitation we will discuss in Section~\\ref{Sec:Discussion} below. The simulations are designed primarily to address the impact of radiative feedback on the assembly of the emerging dwarf galaxy, and on the formation of galactic disks inside it. However, we will also briefly discuss the impact of radiative feedback on the assembly of galaxies in the neighborhood of the simulated dwarf galaxy. Our simulations enable us to provide an improved estimate of the observability of the first galaxies with {\\it JWST}. \\par The organization of this paper is as follows. In Section~\\ref{Sec:Numerics} (as well as in the appendix) we describe our numerical techniques, and in Section~\\ref{Sec:Simulations} we describe the set of simulations that we have carried out. In Sections~\\ref{Sec:Formation} and \\ref{Sec:Reionization} we present the results of our simulations, subsequently discussing the assembly of the dwarf galaxy and the radiative feedback on the IGM and the neighboring galaxies. In Section~\\ref{Sec:Prospects} we use the simulated star formation rates to estimate the observability of the first galaxies with {\\it JWST}. In Section~\\ref{Sec:Discussion} we discuss our results and also address some of the most important limitations. In Section~\\ref{Sec:Summary}, we summarize our work. \\par Throughout this work we assume the $\\Lambda$CDM cosmological model with parameters $\\Omega_{\\rm{m}} = 0.258, \\Omega_{\\rm{b}} = 0.0441, \\Omega_\\Lambda = 0.742, \\sigma_8 = 0.796, n_{\\rm{s}} = 0.963$, and $h = 0.719$, which are consistent with the most recent analysis of the observations with the {\\it Wilkinson Microwave Anisotropy Probe} satellite (\\citealp{Komatsu:2011}). Distances are expressed in physical (i.e., not comoving) units, unless noted otherwise. We will make use of the species number density fractions with respect to hydrogen $\\eta_\\alpha \\equiv n_{\\rm \\alpha} / n_{\\rm H}$, where $\\alpha$ labels the chemical species. \\par ",
        "conclusions": "\\label{Sec:Discussion} \\par The two nested gas disks in our simulations form only after the halo has reached a virial temperature significantly larger than $10^4 \\K$. In light of the scale-free nature of cold dark matter structure formation, this relatively late formation of the disks may be surprising. However, it is consistent with results from previous zoomed simulations of the first atomically cooling galaxies which have not exhibited orderly rotation of the gas (e.g., \\citealp{Wise:2008c}; \\citealp{Greif:2008}; \\citealp{Prieto:2013}). On the other hand, the occurrence of the disks in the dwarf-sized halos fits smoothly in line with simulations of galaxies more massive than the first atomically cooling galaxies assembling at lower redshifts (e.g., \\citealp{Mashchenko:2008}; \\citealp{Pawlik:2011a}; \\citealp{Romano:2011}; \\citealp{Wise:2012}). \\par The results from these previous works may point at a threshold halo mass for the formation of the first disk galaxies in the range $\\sim 10^8-10^{10}\\Msun$ at $z\\sim 10$. This mass scale is similar to the mass scale above which stellar feedback becomes inefficient. However, the fact that simulations of the first atomically cooling galaxies have yielded a turbulent morphology even in the absence of star formation suggests that the two scales are physically unrelated. Indeed, in our simulations, the halo mass at the time of disk formation is insensitive to the inclusion of feedback. Nevertheless, the increased robustness of dwarf galaxies against stellar feedback helps to preserve the disks (e.g., \\citealp{Kaufmann:2007}). We caution that properties other than halo mass such as, e.g., the environment, or the merger history, may be critical to disk formation in the first galaxies (\\citealp{Prieto:2013}).  \\par \\par Our results imply that galaxies with SFRs of $\\sim 0.1 \\Msunyri$ will be among the faintest galaxies {\\it JWST} is likely to detect at $z > 10$, confirming earlier estimates (e.g., \\citealp{Haiman:1998}; \\citealp{Oh:1999}; \\citealp{Tumlinson:2001}; \\citealp{Ricotti:2008}; \\citealp{Zackrisson:2011}; \\citealp{Pawlik:2011a}). According to our simulations, such galaxies reside in halos with masses of $\\sim 10^9 \\Msun$ having stellar masses of $\\sim 10^7\\Msun$. In principle, {\\it JWST} is sufficiently powerful to detect the light from stellar clusters with masses as low as $10^5-10^6 \\Msun$ (e.g., \\citealp{Johnson:2009}; \\citealp{Zackrisson:2011}; \\citealp{Pawlik:2011a}). Our simulations do not support the formation of such massive clusters, which would require local SFRs $\\sim 1 \\Msunyri$, an order of magnitude higher than found here. However, star formation in our simulations is unresolved, and it is possible that star formation in the first galaxies is more clustered or bursty than implied here. In this case, {\\it JWST} could detect galaxies inside halos less massive than considered here, or inside halos that are more strongly affected by feedback than suggested by our simulations. On the other hand, because we do not resolve the formation of stars from first principles, galaxies could be less efficient star-formers than implied by our simulations, and hence be fainter. Such fainter galaxies may still be seen if they are gravitational lensed (e.g., \\citealp{Zackrisson:2012}). \\par Our simulations have ignored a potentially very important physical process, namely the explosion of massive stars in SNe. SNe can provide a strong negative feedback by heating and expelling gas from even relatively massive halos (e.g., \\citealp{MacLow:1999}). Previous works have shown that SN feedback can suppress star formation strongly and lead to bursty star formation histories (e.g., \\citealp{Stinson:2007}). SN feedback further may create a highly spatially inhomogeneous medium, likely enhancing the fraction of low column density sight-lines and hence the fraction of escaping ionizing photons (e.g., \\citealp{Yajima:2009}; but see \\citealp{Dove:2000}). SN feedback may disturb the assembly of disks inside the first galaxies strongly (e.g., \\citealp{Wise:2012}; but see, e.g., \\citealp{Mashchenko:2008}). Moreover, SN feedback may modify the structure of the dark matter halos significantly (e.g., \\citealp{Mashchenko:2008}; \\citealp{Governato:2012}; \\citealp{Brooks:2012}; \\citealp{Garrison:2013}). \\par Our simulations also ignored the chemical enrichment of the gas by the metals synthesized in stars. Simulations that track the production and transport of metals suggest that the transition between metal-free and metal-enriched stellar populations may occur early in the history of the universe (e.g., \\citealp{Tornatore:2007}; \\citealp{Maio:2010}; \\citealp{Wise:2012}). Significant uncertainties, however, remain as to the efficiency of the mixing of metals with the primordial gas and the level of spatial homogeneity of metal enrichment (e.g., \\citealp{Scannapieco:2002}; \\citealp{Ricotti:2008}). This, together with the fact that not all stars are expected to explode in SNe and enrich the gas but may instead collapse directly into black holes (e.g., \\citealp{Heger:2003}), leaves open the possibility of the formation of metal-free stars in select regions of the universe down to relatively low redshifts (e.g., \\citealp{Tornatore:2007}; \\citealp{Trenti:2009}; \\citealp{Johnson:2010}; \\citealp{Fumagalli:2011}; \\citealp{Simcoe:2012}). However, even if all stars would collapse into black holes without SNe, feedback from accretion onto the black holes may still affect the evolution of the galaxies in a manner not captured by our simulations (e.g., \\citealp{Ricotti:2004}; \\citealp{Alvarez:2009}; \\citealp{Jeon:2012}). \\par  We have presented cosmological smoothed particle hydrodynamics simulations of a dwarf galaxy assembling in a halo reaching $10^9 \\Msun$ at $z = 10$. The simulations were identical to our earlier simulations of such a galaxy in that they followed the non-equilibrium chemistry and cooling of primordial gas. They improved on our earlier simulations by including the formation of massive metal-free stars. To investigate the radiative feedback from these stars, we compared a simulation in which star particles emitted both molecular hydrogen dissociating and hydrogen/helium ionizing radiation and a simulation in which star particles emitted only dissociating radiation with a simulation inside which star particles remained dark. \\par Our main results are: \\begin{itemize}  \\item Dissociating and ionizing radiation exert a strong negative feedback by suppressing star formation in the main minihalo progenitor of the dwarf galaxy, but have little effect on star formation as soon as the progenitor evolves into an atomically cooling galaxy.  \\item Radiative feedback suppresses the central dark matter densities in the dwarf galaxy main progenitor minihalo relative to the densities found in the simulation without radiation. The dark matter density profile of the dwarf galaxy is singular isothermal independent of the inclusion of radiation shortly after the minihalo has evolved into an atomic cooling halo.  \\item The dwarf galaxy halo hosts two nested disks below $z \\lesssim 12.5$. The formation history and structure of the disks are insensitive to the inclusion of dissociating and ionizing radiation. These results are consistent with a picture in which the first disk galaxies form inside dark matter halos with masses $\\gtrsim 10^9 \\Msun$ at $z \\gtrsim 10$.  \\item The inclusion of dissociating and ionizing radiation lowers the baryon fractions inside the minihalos in the neighborhood of the dwarf galaxy. The baryon fractions are lowest in minihalos with masses $\\lesssim 10^6 \\Msun$, a consequence of Jeans filtering and photoevaporation from external ionizing sources, and in minihalos with masses $\\sim 10^7-10^8 \\Msun$, here primarily a consequence of photoevaporation of gas by internal ionizing sources.  \\item Galaxies with star formation rates $\\sim 0.1 \\Msunyri $ will be among the faintest galaxies the upcoming {\\it James Webb Space Telescope} will detect in deep exposures of the $z \\gtrsim 10$ universe. Our simulations suggest that such galaxies reside in halos with masses $\\sim 10^9 \\Msun$ and have stellar masses $\\sim 10^7 \\Msun$.  \\end{itemize} \\par We caution that our conclusions are subject to statistical uncertainties implied by the small volume of the high-resolution region in our simulations. Another major shortcoming of our simulations is the lack of feedback from supernova explosions. Such feedback can potentially have a significant impact on the evolution of low-mass galaxies. Feedback from supernovae may heavily disturb the assembly of disks, and strongly decrease the star formation rates inside dwarf galaxies, thus affecting also estimates of their observability. Our simulations also did not account for the chemical enrichment of the interstellar and intergalactic gas and the associated transition from metal-free to metal-enriched stars. The effects of supernova feedback and chemical enrichment are left to be investigated in future work."
    },
    "1208/1208.1524_arXiv.txt": {
        "abstract": "We present observations and magnetic field modeling of the large polar crown prominence that erupted on 2010 December 6. Combination of SDO/AIA and STEREO$\\_$Behind/EUVI allows us to see the fine structures of this prominence both at the limb and on the disk. We focus on the structures and dynamics of this prominence before the eruption. This prominence contains two parts: an active region part containing mainly horizontal threads, and a quiet Sun part containing mainly vertical threads. On the northern side of the prominence channel, both AIA and EUVI observe bright features which appear to be the lower legs of loops that go above then join in the filament. Filament materials are observed to frequently eject horizontally from the active region part to the quiet Sun part. This ejection results in the formation of a dense-column structure (concentration of dark vertical threads) near the border between the active region and the quiet Sun. Using the flux rope insertion method, we create non-linear force-free field models based on SDO/HMI line-of-sight magnetograms. A key feature of these models is that the flux rope has connections with the surroundings photosphere, so its axial flux varies along the filament path. The height and location of the dips of field lines in our models roughly replicate those of the observed prominence.  Comparison between model and observations suggests that the bright features on the northern side of the channel are the lower legs of the field lines that turn into the flux rope. We suggest that plasma may be injected into the prominence along these field lines. Although the models fit the observations quiet well, there are also some interesting differences. For example, the models do not reproduce the observed vertical threads and cannot explain the formation of the dense-column structure. ",
        "introduction": "Solar prominences are relatively cool structures embedded in the million-degree corona \\citep{1985SoPh..100..415H, 2010SSRv..151..243L, 2010SSRv..151..333M}. In H$\\alpha$ when viewed above the solar limb, prominences appear as bright structures against the dark background, but when viewed as ``filaments'' on the solar disk they are darker than their surroundings. We will use the terms ``filament\" and ``prominence\" interchangeably in general. A filament is formed in a filament channel, which is defined as a region in the chromosphere surrounding a polarity inversion line (PIL) where the chromospheric H$\\alpha$ fibrils are aligned with the PIL \\citep{1971SoPh...19...59F,1998ASPC..150..257G}. Filaments can be found inside activity nest consisting of multiple bipolar pairs of spots (``active region filaments\"), at the border of active regions (``intermediate filaments\"), and on the quiet Sun (``quiescent filaments\"), including the polar crown.  Filaments typically consist of  three structural components: a spine, barbs, and two extreme ends \\citep{1998SoPh..182..107M, 2008ASPC..383..235L, 2011SSRv..158..237L}. The spine runs horizontally along the top of the filament, although there may be sections along the filament where the spine is nearly invisible. The barbs protrude from the side of the filament and when observed closer to the limb, the barbs are seen to extend down from the spine to the chromosphere below. The ends, also called ``legs\", may be a collection of threads that appear to terminate at a single point or at multiple points. When a quiescent filament is viewed with high resolution on the solar disk, H$\\alpha$ observations indicate that each of these three structural components consist of thin thread-like structures \\citep{2008ASPC..383..235L, 2011SSRv..158..237L}. In active regions the thin filament threads often appear to be mostly horizontal \\citep{2007Sci...318.1577O, 2008ASPC..383..235L}, while quasi-vertical threads (``hedgerows\") are often seen in the quiescent prominences \\citep{2008ApJ...676L..89B, 2010ApJ...716.1288B, 2008ApJ...689L..73C, 2008ApJ...686.1383H}.  Prominence plasma is highly dynamic, exhibiting horizontal \\citep{2008ApJ...689L..73C} or vertical \\citep{2008ApJ...676L..89B}  flows. The flows reported so far are either unidirectional \\citep {2008ApJ...676L..89B, 2008ApJ...689L..73C} or bi-directional  \\citep{1998Natur.396..440Z, 2003SoPh..216..109L}. The latter is known as counter-streaming.  The magnetic field plays a primary role in filament formation, stability, and eruption \\citep{Priest1989, Tandberg-Hanssen1995, 2010SSRv..151..333M}. However, the magnetic structure of prominences is still not fully understood, with many observations and theoretical models differing on the exact nature of the magnetic field. Various models for prominence magnetic structure can be summarized as follows. In most of the models, filament plasmas are assumed to be located near the dips of magnetic field lines. These models can be classified as ``sheared arcade model'' \\citep{1994ApJ...420L..41A, 2000ApJ...539..954D, 2002ApJ...567L..97A} and ``flux rope model\", \\citep{1974A&A....31..189K, 1983SoPh...88..219P, 1989ApJ...344.1010P, 1989ApJ...343..971V, 1994SoPh..155...69R, 1995ApJ...443..818L, 1998A&A...335..309A, 2001ApJ...560..476C, 2004ApJ...612..519V, 2006ApJ...641..590G, 2008SoPh..248...29D}. In these models, the prominence plasmas are embedded in a  large sheared arcade or a flux rope that lie horizontally above the PIL. Flux rope models may be split into two categories: weakly twisted flux rope and highly twisted flux rope. Weakly twisted flux rope models are similar to sheared arcades with one key difference as argued by \\citet{2010SSRv..151..333M}. For flux rope models the flux rope and overlying arcade are independent flux systems with a separatrix surface between them. In contrast, for a sheared arcade model only a single flux system exists. In addition to the above mentioned 3D global models, there are also two other dip models which focus on the local support of vertical threads in hedgerow prominence, i.e., ``sagged horizontal fields model\" \\citep{2008ApJ...689L..73C, 2010ApJ...714..618C} and ``tangled magnetic field model\" \\citep{2010ApJ...711..164V}. In the earlier model, the vertical threads are stacks of plasma supported against gravity by the sagging of initially horizontal magnetic field lines. In the latter model, tangled fields exist in a vertical current sheet of small width that is confined by the vertical fields on either side of the sheet. Neither of these models describes the global 3D topology of the magnetic field supporting the filament. In contrast to models mentioned above, the ``field aligned thread model\" developed by \\citet{2008ASPC..383..235L, 2011SSRv..158..237L} does not contain dips. This model is a 3D empirical magnetic model based on high-resolution H$\\alpha$ observations and is based on the assumption that the observed fine scale structures are parallel to the magnetic field. In this model, the filament plasma is located on magnetic arches that are highly sheared in the direction along the PIL. Some field lines run along the entire length of the filament and outline the filament ``spine\". Other shorter ones run partially along the spine, but spread out from it and connect down to minority polarity elements on either side of the PIL. These shorter structures represent the filament barbs. In this case, the plasma is supported by MHD waves. However, relatively high frequencies and wave amplitudes are required, and it is unclear why such waves would not lead to strong heating of the prominence plasma.  The purpose of the present paper is to analyze observations of a quiescent prominence, and to develop 3D models of its magnetic structure. Filaments in active regions have been modeled extensively in the past. Several authors have developed non-linear force-free field (NLFFF) models which are based on magnetic observations. For example, \\citet{2007A&A...468..701R} construct NLFFF models by ``extrapolating\" observed photospheric vector fields into the corona \\citep[also see][]{2002A&A...392.1119R, 2010ApJ...715.1566C}. van Ballegooijen (2004) developed an alternative method for constructing NLFFF models of filament flux ropes. The method involves inserting a flux rope into a potential field based on an observed photospheric magnetogram, and then evolving the field in time to an equilibrium state using magneto-frictional relaxation. This method has been used by \\citet{2008ApJ...672.1209B} and \\citet{2009ApJ...704..341S, 2009ApJ...691..105S, 2011ApJ...734...53S, 2012ApJ...746...81A} to study active regions with filaments, and by \\citet{2009ApJ...703.1766S, 2012ApJ...744...78S, 2012ApJ...750...15S} to study the evolution of soft X-ray sigmoid. The papers mentioned above suggest that the observed active region filaments can be well explained by a flux rope model. In this paper we apply this flux rope insertion method to a quiescent prominence for the first time. ",
        "conclusions": "A large polar crown prominence partially erupted on 2010 December 6. To understand the magnetic support of the prominence, we focus on the structure and dynamics of the prominence before and after the eruption. The 6-day SDO/AIA observations of the filament near the east limb before the eruption suggest that this filament can be divided into 3 parts. The active region part near the west end contains mainly horizontal threads. The middle part near the quiet Sun contains vertical threads with overlying horizontal threads. While the east end of the prominence that survived the eruption appears to be composed of mainly vertical threads. The post-eruption limb observations by AIA suggest that the prominence left behind contains mainly vertical threads, while a regular thin dark filament structure appears in the on-disk observations by EUVI. This thin dark filament structure may be the horizontal filament threads located on top of the vertical threads. Another possibility is that it is just the accumulation of the lined up dark vertical threads, and there are no horizontal threads. Such structure, consisting of horizontally aligned threads, was reported by \\citet{2004SoPh..221..297S} for another filament. Corresponding H$\\alpha$ appearance of such threads was discussed by \\citet{2006ApJ...643L..65H}.  STEREO/EUVI observes straight and faint features representing the lower legs of the overlying coronal arcades on the two sides of the filament channel. On the quiet Sun part, bright features with no clear counterparts are also identified on the northern side of the channel. This emission asymmetry is consistent with the existence of loops that originates from the northern side of the channel, then arch over the filament and disappear or join in the filament as shown in earlier limb observations at 171~\\AA~by AIA. A clear cavity corresponding to the filament channel is observed on December 4 at 335~\\AA~by AIA and XRT. The observations suggests that the filament is located on the southern edge of the cavity rather than at the center. The bright features (lower legs of the coronal arcades) on the north side of the cavity are brighter than those on the south side. The skew of the post-eruption arcade suggests that this filament is sinistral.  The prominence before the eruption is very dynamic. AIA observes horizontal counter-streaming overlying the vertical threads. STEREO$\\_$B/EUVI observes filament material ejection from the active region part to the quiet Sun part. This ejection leads to strong oscillations of the vertical threads and the formation of a dense column near the border between the active region and quiet Sun. The dense column refers to strong accumulation of dark vertical filament threads.  Using the flux rope insertion method developed by \\citet{2004ApJ...612..519V} we construct two non-linear force-free field models. The poloidal flux ($1\\times 10^{20}$ Mx cm$^{-1}$) in Model 1 is half of that in Model 2, so the flux rope in Model 1 is less twisted. In both models, the axial fluxes are $8 \\times 10^{20}$ Mx  and $2 \\times 10^{20}$ Mx in the active region and quiet Sun, respectively. Since the prominence is observed at the east limb, the models are constructed based on the combined LOS photospheric magnetogram observed by SDO/HMI several day after the prominence eruption. Both models lie at the edge of the stable regime.  Our models match some of the observed features quiet well. For example, the height and location of the field line dips in our models appear to match those of the prominence threads, though the narrow vertical thread-like structure is not reproduced in our model. The lower part of the magnetic field lines from our models can replicate the bright features on the two side of the filament channel quite well. Our model also confirms our previous interpretation of the emission asymmetry on the two sides of the channel on the quiet Sun \\citep{2010ApJ...721..901S}. This emission asymmetry is due to the fact that some of the bright features on the northern side (the brighter side) of the channel represent the field lines that turn into the flux rope. Therefore, no counterparts of these features can be found on the southern side of the channel. We suggest that plasma may be injected into the prominence along the aforementioned field lines that turn into the flux rope, since these field lines start from the photosphere. This is consistent with the thermal nonequilibrium mechanism for prominence formation as suggested by Antiochos, Karpen, and colleagues \\citep[e.g., ][]{1991ApJ...378..372A, 2005ApJ...635.1319K, 2012ApJ...746...30L}. In this model, localized heating above the flux tube footpoints produces evaporation of the chromospheric plasma, which condenses in the coronal part of the flux tube.   Acknowledgments: We thank the anonymous referee for valuable comments to improve this paper. Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and the NSC (Norway). We thank the team of SDO/AIA, SDO/HMI, STEREO/EUVI, Hinode/XRT, KSO, and SOLIS for providing the valuable data. The EUVI and HMI data are downloaded via the Virtual Solar Observatory and the Joint Science Operations Center. Y. Su acknowledges Dr. Mark Weber for helpful discussions and thanks Dr. Suli Ma for helping on the radial filter technique. This project is partially supported under contract NNM07AB07C from NASA to the Smithsonian Astrophysical Observatory (SAO) and SP02H1701R from LMSAL to SAO as well as NASA grant NNX12AI30G."
    },
    "1208/1208.6132_arXiv.txt": {
        "abstract": "The first complete submillimetre spectrum (190-670$\\mu$m) of the Seyfert 2 galaxy NGC\\,1068 has been observed with the SPIRE Fourier Transform Spectrometer onboard the {\\it Herschel} Space Observatory. The sequence of CO lines (J$_{up}$=4-13), lines from \\water, the fundamental rotational transition of HF,  two o-H$_2$O$^{+}$ lines and one line each from CH$^{+}$ and OH$^{+}$ have been detected, together with the two [CI] lines and the [NII]205$\\mu$m line. The observations in both single pointing mode with sparse image sampling and in mapping mode with full image sampling allow us to disentangle two molecular emission components, one due to the compact circum-nuclear disk (CND) and one from the extended region encompassing the star forming ring (SF-ring). Radiative transfer models show that the two CO components are characterized by density of $n({\\rm H_2})$=$10^{4.5}$ and $10^{2.9}$ cm$^{-3}$ and temperature of $T_{\\rm kin}$ =100K and  127K, respectively. The comparison of the CO line intensities with photodissociation region (PDR) and X-ray dominated region (XDR) models, together with other observational constraints, such as the observed CO surface brightness and the radiation field, indicate that the best explanation for the CO excitation of the CND is an XDR with density of n(H$_2$) $\\sim$ 10$^4$ cm$^{-3}$ and X-ray flux of 9 erg~s$^{-1}$~cm$^{-2}$, consistent with illumination by the active galactic nucleus, while the CO lines in the SF-ring are better modeled by a PDR. The detected water transitions, together with those observed with the \\her~ PACS Spectrometer, can be modeled by an LVG model with low temperature ($T_{\\rm kin}\\sim$ 40K) and high density ($n({\\rm H_2})$ in the range $10^{6.7}$ --- $10^{7.9}$ cm$^{-3}$). The emission of \\hdop\\ and \\ohp\\ are in agreement with PDR models with cosmic ray ionization.  The diffuse ionized atomic component observed through the [NII]205$\\mu$m line is consistent with previous photoionization models of the starburst. ",
        "introduction": "\\label{intro}  NGC\\,1068 (Messier 77) is a nearby ($cz$ = 1137 km s$^{-1}$) and bright (L$_{IR}$=L$_{8-1000\\mu m}$ $\\sim$2 $\\times$ 10$^{11}$ L$_{\\odot}$, \\citealt{bla97}) Seyfert galaxy, often considered as the prototypical Seyfert type 2 galaxy. However, since the discovery of the broad permitted lines in the polarized optical spectrum of NGC\\,1068 \\citep{anto85}, it has become clear that this galaxy was in reality a {\\it hidden broad line region} galaxy. The general distinction between the two types of Seyfert galaxies might be due only to orientation effects, according to the so called Unification model \\citep{anto93}.  Being the strongest nearby Seyfert 2 galaxy, it has been observed extensively over the whole electromagnetic spectrum. Molecular (CO and HCN) observations have shown a prominent starburst ring (hereafter SF-ring) at a radius of 1.0-1.5 kpc and a central circum-nuclear disk (hereafter CND) with a diameter of $d$ $\\sim$ 300 pc \\citep{tac94,sch00}. Near-IR observations \\citep{sco88,thr89} have clearly revealed a 2.3 kpc stellar bar. The compact ($\\sim$1 pc) hot dust source in the nucleus of NGC\\,1068, measured using near-infrared speckle imaging and integral field spectroscopy \\citep{tha97}, is probably heated by the AGN's strong radiation field and possibly associated with the postulated dense circum-nuclear torus. A jet was observed from centimeter to millimeter wavelengths extending out to several kiloparsecs from the center \\citep[e.g.,][]{kri06,gal04}. Mid-IR observations revealed hot and ionized gas biconically following the path of the radio jet \\citep[e.g.,][and references therein]{mul09} and indicating the existence of a parsec-scale warm dust torus \\citep{jaf04}. Recent interferometric observations of CO (3-2) and CO (1-0) show that the CO (3-2) emission peaks in the central region within $\\sim$ 5 $\\arcsec$ from the nucleus, while the CO(1-0) emission is mainly located along the spiral arms \\citep{tsa12}.  Spectroscopic coverage of NGC\\,1068  has been extensive. In the mid-IR to far-IR wavelength range, spectra have been measured by the ISO \\citep{kes96} SWS \\citep{deg96}  and LWS  \\citep{cle96} spectrometers \\citep[][respectively]{lut00,spi05}, covering the 2.4-45~$\\mu$m and 43-197~$\\mu$m spectral ranges. At millimeter wavelengths it was recently observed from the ground in the 190-307 GHz (976-1578~$\\mu$m) range \\citep{kam11}. The submillimeter waveband is one of the few spectral regions that have not been so far explored; the new observations made by the Spectral and Photometric Imaging Receiver (SPIRE) \\citep{gri10} Fourier Transform Spectrometer (FTS) \\citep{nay10a}, onboard the \\her~Space Observatory \\citep{pil10}, covering the spectral range from 190~$\\mu$m to 670~$\\mu$m, fill most of this gap with the first complete submillimeter spectrum of NGC\\,1068.  Submillimeter  spectral measurements are of particular interest as NGC\\,1068 is a prime candidate to study the effects of the AGN onto the circum-nuclear material and the surrounding disk. In particular, the study of the excitation conditions and chemistry of the CND appear already from the existing ground-based molecular line observations to be very different from the starburst galaxies environments. The peculiar line ratios of different molecular transitions, mostly HCN, HCO$^+$, and $^{12}$CO, led to the suggestion that the CND of NGC\\,1068 harbors a giant X-ray-dominated region \\citep[XDR, e.g.,][]{rot91, mal96, use04, koh08}. The HCN and HCO$^+$ molecular line studies of \\citet{kri08,kri11} confirm an increased abundance of HCN and/or increased kinetic temperatures. CO lines can also discriminate between  ``classical'' photodissociation regions (PDRs) and X-ray dominated regions (XDRs) \\citep[e.g.][]{mei05}. Using the intermediate J rotational lines from the CO molecule, from J$_{up}$=4 to J$_{up}$=13, which can be observed with the SPIRE FTS, we want to test if indeed in NGC\\,1068 an XDR is needed to explain the spectral line energy distribution originating from the CND. The case of the ultraluminous IR galaxy Mrk231 already demonstrated that the \\her -SPIRE data are indeed able to discriminate between the two emission mechanisms and therefore detect the effects of the AGN \\citep{wer10}. \\her -PACS observations of the high-J CO lines (J$_{up} \\ge 14$) detected in NGC\\,1068 have been recently presented by \\citet{hai12}. The two components, at high and medium excitation, needed to explain the observed CO lines arising from the central 10$\\arcsec$ region from J$_{up}$= 14 to J$_{up}$= 24 can be excited by X-ray or shock heating, while far-UV heating is unlikely.  The \\her -SPIRE spectroscopic observations of NGC\\,1068 presented here have been collected under the guaranteed time key project ``Physical Processes in the Interstellar Medium of Very Nearby Galaxies'' (PI: Christine Wilson). Within the same observational program, SPIRE and Photodetector Array Camera and Spectrometer (PACS) \\citep{pog10} photometric images have been collected, which will be presented in a forthcoming paper with the detailed analysis of the continuum emission (Spinoglio et al 2012, in prep.). ",
        "conclusions": "We summarize here the results of this work. The first complete submillimetre (190--670~$\\mu$m) spectrum of the Seyfert type 2 galaxy NGC\\,1068 reveals the full sequence of CO pure rotational lines from J$_{up}$=4 to J$_{up}$=13. The radiation transfer analysis of these lines shows the presence of two physically distinct components: the first one originating from the circum-nuclear disk (CND) of few arcseconds in diameter ($\\sim$ 4$\\arcsec$) and the second one excited in the star forming ring (SF-ring) with a diameter ten times larger ($\\sim$ 40$\\arcsec$). These results indicate a kinetic temperature of CO  of $T_{\\rm kin}$=100~K and  127~K, a gas density of $n({\\rm H_2})$=$10^{4.5}$ and $10^{2.9}$ cm$^{-3}$ and a derived molecular hydrogen mass of M(H$_{2}$) $\\sim$ 2.4$\\times$10$^{7}$~M$_{\\sun}$ and M(H$_{2}$) $\\sim$ 3.5$\\times$10$^{8}$~M$_{\\sun}$, for the compact and extended regions, respectively.  The comparison of the observed CO line intensities with predictions of photodissociation (PDR) and X-ray dominated regions (XDR) models shows that the circum-nuclear disk emission can be modeled equally well by both types of models, while the CO lines in the star-forming ring can be modeled by a photodissociation region only. However some observational constraints, such as the total CO surface brightness and the required radiation field, indicate that the most plausible explanation for the CO excitation of the CND is an XDR with density of n(H$_2$) $\\sim$ 10$^4$ cm$^{-3}$ and X-ray flux of 9 erg~s$^{-1}$~cm$^{-2}$, consistent with the AGN illumination. In contrast, the excitation of the SF-ring component is due to PDR emission originating from the young stars/HII regions in the spiral arms.  The water lines that we have detected with SPIRE, together with those observed by PACS (S. Hailey-Dunsheath 2012, private comm.), have been modeled with an LVG model to constrain the physical conditions of the water excitation. We have found that the kinetic temperature is T$_{kin}$=40~K, the molecular hydrogen density is $n(H_2)=4 \\times 10^6 - 8\\times 10^7$ and the column density is of  order N(H$_2$O)=8 $\\times 10^{17}$ cm$^{-2}$ for both water forms.  The computed column densities of the molecular ions detected (\\hdop\\ and \\ohp\\ ) are in agreement with PDR models that include cosmic ray ionization.  The fundamental rotational transition of HF has been detected in emission in NGC\\,1068 and we infer a column density of N(HF)$\\sim 7.5\\times 10^{12}$ cm$^{-2}$.  For the two [CI] transitions, we derived a kinetic temperature of 22.5$\\pm$2.4 K in LTE approximation, which is much lower than the temperatures traced from the intermediate-J CO molecular gas.  The molecular masses that we derived from our analysis are in good agreement with the masses estimated from both CO interferometric measurements of low-J lines and mid-infrared H$_2$ emission lines.  Finally we show that  the intermediate-J CO and [CI] lines in galaxies with L$\\sim10^{12}L_{\\sun}$ can be observed from planned and future ground-based and space telescopes up to redshift of z $\\sim$ 0.5, making their diagnostic power an important tool to study galaxy evolution at intermediate redshift."
    },
    "1208/1208.6304_arXiv.txt": {
        "abstract": "We present observational evidence for the inhibition of bar formation in dispersion-dominated (dynamically hot) galaxies by studying the relationship between galactic structure and host galaxy kinematics in a sample of 257 galaxies between 0.1 $<$ z $\\leq$ 0.84 from the All-Wavelength Extended Groth Strip International Survey (AEGIS) and the Deep Extragalactic Evolutionary Probe 2 (DEEP2) survey.  We find that bars are preferentially found in galaxies that are massive and dynamically cold (rotation-dominated) and on the stellar Tully-Fisher relationship, as is the case for barred spirals in the local Universe.  The data provide at least one explanation for the steep ($\\times$3) decline in the overall bar fraction from z=0 to z=0.84 in L$^*$ and brighter disks seen in previous studies.  The decline in the bar fraction at high redshift is almost exclusively in the lower mass (10 $<$ log M$_{*}$(\\Msun)$<$ 11), later-type and bluer galaxies. A proposed explanation for this  ``downsizing'' of the bar formation / stellar structure formation is that the lower mass galaxies may not form bars because they could be dynamically hotter than more massive systems from the increased turbulence of accreting gas,  elevated star formation, and/or increased interaction/merger rate at higher redshifts.  The evidence presented here provides observational support for this hypothesis.  However, the data also show that not every disk galaxy that is massive and cold has a stellar bar, suggesting that mass and dynamic coldness of a disk are necessary but not sufficient conditions for bar formation -- a secondary process, perhaps the interaction history between the dark matter halo and the baryonic matter,  may play an important role in bar formation. ",
        "introduction": "\\label{intro}  The presence of galactic structures such as bars is as an  important signpost in the evolution of a galaxy disk.  Analystical work and simulations have shown that, once a galaxy disk is sufficiently massive and dynamically cold, the formation of a stellar bar is relatively fast ($\\sim$hundred million years) (e.g., \\citealt{hohl71, kalnajs72, ostriker73, sellwood93, athanassoula02, athanassoula03, heller07}).  But bar formation can be delayed either by an {\\em initially} dominant dark matter (DM) halo and/or a dynamically hot/dispersion-dominated disk \\footnote{We use the terms dynamically hot and dispersion-dominated interchangeably throughout the paper.} \\citep{athanassoula86}\\   \\\\  An {\\em initially} dominant DM halo strongly impacts the time scale for bar formation delaying the onset of the bar instability \\citep{athanassoula02}.  The bar that ultimately forms in such a system is stronger than a bar that would form in an otherwise non-DM dominated galaxy because the DM  halo acts as an efficient sink of angular momentum and energy for baryons, which are redistributed to form the bar.  As the bar grows it pushes material inwards so that the baryonic matter can become the dominant mass component in the inner parts of galaxies \\citep{athanassoula02b}. Simulations also show that a dynamically hot disk delays bar formation \\citep{athanassoula86, athanassoula03} because when random motions of stars in a disk have a higher amplitude than rotational ordered motions, the bar instability cannot grow quickly.  This may even push the bar formation time scale beyond a Hubble time.  \\\\ %    A recent COSMOS study of over two thousand L$^*$ and brighter, face-on ($i < $ 65$^{\\circ}$) disk galaxies showed that the overall bar fraction (f$_{bar} $ = total number of barred galaxies divided by the total number of disk galaxies) in disk galaxies declines sharply from f$_{bar} \\sim$ 0.65 at z=0  to f$_{bar}$ $<$ 0.2 at z=0.84 \\citep{sheth08}.  It is crucial to note that the COSMOS sample is a complete sample only for disks with stellar masses, M$_* > $ 10$^{10}$\\Msun; the published results of the bar fraction evolution apply {\\em only} to this mass range \\citep{sheth08, cameron10}.  Therefore, studies with samples of lower mass galaxies (M$_* < $ 10$^{10}$\\Msun, such as those typically done for nearby galaxies (e.g.,  using SDSS)) are not directly comparable to high redshift studies.  The evolution of the bar fraction with redshift is not uniform across all disk galaxies.  As a function of redshift, f$_{bar}$ is strongly correlated  with the host galaxy mass, color and bulge dominance (see Figures 2--5 in \\citet{sheth08}).  The most massive stellar disks (M$_* \\ge $10$^{11}$\\Msun), which are also redder and have a larger bulge, already had f$_{bar} >$ 0.5 at z$\\sim$0.8, nearly their present-day value of their bar fraction.  In sharp contrast, the lower stellar mass systems (M$\\sim$10$^{10}$ \\Msun),  had f$_{bar} \\ll $ 0.2 at z$\\sim$0.8.  Over the last 7 Gyr, the lower mass galaxies have evolved the fastest, increasing their bar fraction by more than a factor of three, to their present day value of f$_{bar} \\sim$ 0.65.  This behavior is another form of ''downsizing\" \\citep{cowie96}.  The dynamics of high redshift disks has been a hot topic of study in recent years (e.g., \\citealt{kassin07, schreiber09,cresci09,lorenzo09,davies11,miller11, kassin12}).  At high redshifts (z$\\ge$2) there is evidence for both rotation- and  dispersion-dominated disks (e.g., \\citep{law09, cresci09,wright09}), although the evolution of the disk kinematics and assembly is not well-understood. The dynamics of a galaxy  must change as it acquires mass, undergoes interactions/mergers and forms stars.  In this paper we seek to understand how the disk dynamics are influencing the formation of bars.   \\\\  In a 2007 study of over $\\sim$500 galaxies from 0.2$<$z$<$1.2, \\citet{kassin07} found that major-mergers,  disturbed and compact systems are preferentially off the stellar mass Tully-Fisher (TF) relationship towards lower rotational velocities.  In contrast for the local Universe, \\citet{barton01}, who examined 90 close pairs, found that only eight scattered off the TF relationship.  \\citep{kannappan02} analyzed the residuals in the TF for a wide variety of galaxy morphologies and environments from the Nearby Field Galaxy Survey and found that the scatter in the TF did not change once corrections for dust extinction and star formation were applied, although the scatter did increase for non-spiral galaxies.  They also found that dwarf galaxies did not follow the TF with dwarfs scattering on both sides of the TF (see \\citealt{kannappan02} for an in-depth discussion).  At high redshifts,  Figure 1 of \\citet{kassin07} suggests that more early type spirals (blue squares in their Figure 1) are on the classical stellar TF relationship compared to late-type/irregular spirals.  It appears that over time, more and more of late-type/irregular galaxies arrive onto  the stellar-TF.   A different study of disk-like galaxies by \\citet{miller11} has argued that there is no significant evolution in the stellar mass TF relationship to z$\\sim$1, although there is an evolutionary trend in the B-band TF.  While the precise evolution of the stellar TF is not known, we make use of the existing measurements of disk properties (mass, rotational velocity and velocity dispersion) from \\citet{kassin07} and compare these for different types of galaxies. ",
        "conclusions": "The main result of this paper is that bars are not present in dispersion-dominated disk galaxies.  The data suggest an evolutionary sequence in the assembly of disks and formation of the familiar galactic structures such as bars that we see today.   The clump-cluster and chain galaxies are believed to be an early phase of present-day spiral galaxies undergoing a burst of star formation in large, gravitationally unstable clumps in a cold, gaseous disk (e.g., \\citealt{elmegreen05,elmegreen09}).  As these disks evolve and accrete more cold material from the large scale structure filaments, they should evolve towards more rotationally-supported disks, the type that we see on the Tully-Fisher relationship (see also discussion in \\citealt{kassin07}).  While it is not clear how these systems migrate from the left hand side of the diagram to the right, the expected evolution is likely to be towards higher rotational velocities and stellar masses (up and to the right) on Figures \\ref{tf}.   The data also suggest that bars may be growing from short to long as the disks evolve to colder and more massive systems, indicating that bar-driven heating of the disk is less significant than the competing cooling processes.  \\begin{figure}[ht!] \\centering \\begin{tabular}{ll} \\plotone{f9.eps} \\end{tabular} \\caption{Redshift distribution for the 257 galaxies in this sample is shown here.  The median redshift of the sample is 0.46.  There are too few galaxies in each redshift bin to discuss evolutionary trends with both mass and redshift which will have to wait for even larger surveys.} \\label{histz} \\end{figure}  Previous studies have shown a strong correlation between the bar fraction, stellar mass of the galaxy and redshift such that massive galaxies ($>$ 10$^{11}$ \\Msun) had a high ($>$50\\%) bar fraction at z$\\sim$0.85, whereas lower mass galaxies (10$^{10}$\\Msun) had a bar fraction $<$ 20\\% \\citep{sheth08}. The evolution of the bar fraction was differential over the last 7 Gyr with the fastest growth of the bar fraction occurring in the low mass, blue, late type spirals of masses between 10$^{10}$ and 10$^{11}$ \\Msun \\citep{sheth08,cameron10}.  The present DEEP2/AEGIS sample is too small to measure the redshift and mass dependent evolution of these galaxies.  The redshift distribution is shown in Figure \\ref{histz}. While we do see a segregation along the stellar mass axis between compact galaxies and disk (barred and unbarred) galaxies there are too few galaxies to infer any trends with mass and redshift - further analysis with larger data sets will be very useful to interpret the evolutionary trends with mass and redshift.  Finally the compact systems, which are primarily dispersion-dominated, are seen over the entire redshift range of this survey and are therefore not necessarily only exotic high redshift systems, as was found in \\citet{kassin07}.  The fate of these objects is another interesting area of study, especially at lower redshifts, where high spatial resolution (and higher signal to noise) are available.  Although our data shed some light on the conditions that delay bar formation, the large number of non-barred galaxies that are massive, cold and rotationally supported remains a mystery.  \\citet{courteau03} have already shown that there is no obvious difference in the placement of barred and unbarred spirals on the TF in the local Universe.  Locally as many as 30-35\\% of the disk galaxies are unbarred.  And so we conclude that while dynamic coldness and sufficient stellar mass are necessary conditions for the formation of a bar, they are not sufficient. Mergers and interactions are other processes  which could play a role in bar formation but their impact is difficult to quantify because they can create long-lived or  transient bars or they can destroy existing bars \\citep{gerin90,barnes91,mihos94,mihos96,romanodiaz08}).   Finally, another important process for bar formation is the interaction history between the baryonic matter and the dark matter halo, especially in the inner parts of galaxies because the dark matter halo can act as sink of angular momentum and energy for the baryonic matter settling into the central bar (e.g., \\citet{athanassoula02})."
    },
    "1208/1208.1706_arXiv.txt": {
        "abstract": "Sunspot numbers form a comprehensive, long-duration proxy of solar activity and have been used numerous times to empirically investigate the properties of the solar cycle. A number of correlations have been discovered over the 24 cycles for which observational records are available. Here we carry out a sophisticated statistical analysis of the sunspot record that reaffirms these correlations, and sets up an empirical predictive framework for future cycles. An advantage of our approach is that it allows for rigorous assessment of both the statistical significance of various cycle features and the uncertainty associated with predictions. We summarize the data into three {\\sl sequential} relations that estimate the amplitude, duration, and time of rise to maximum for any cycle, given the values from the previous cycle. We find that there is no indication of a persistence in predictive power beyond one cycle, and conclude that the dynamo does not retain memory beyond one cycle. Based on sunspot records up to October 2011, we obtain, for Cycle 24, an estimated maximum smoothed monthly sunspot number of 97 $\\pm$ 15, to occur in January--February 2014 $\\pm$ 6 months. ",
        "introduction": " ",
        "conclusions": "We now discuss the results from the fitted hierarchical model. This is a more coherent analysis than the separate fitting of the two stages as described in Section~\\ref{sec:statmodel}.  Thus, the results in this section represent our final estimates.  \\subsection{Cycle-to-Cycle Dependencies} \\label{sec:cyccor}  Of particular interest are the estimates of the three second-stage relationships. These estimates are summarized in Table~\\ref{tbl:result}, based on an MCMC calculation that fits the two stages jointly, using available data from January 1749 to October 2011. As noted before, because of the large number of observations, cycle characteristics such as $c_i, t_0^{(i)}, t_{\\rm max}^{(i)}$, and $t_1^{(i)}$ are well constrained by the first-stage model (Equation~(\\ref{eq:stage1})), and hence the estimates reported in Table~\\ref{tbl:result} are quite close to those based on an ordinary least squares (OLS) regression using fixed $c_i, t_0^{(i)}, t_{\\rm max}^{(i)}, t_1^{(i)}$. For example, with $c_i, t_0^{(i)}, t_{\\rm max}^{(i)}$ fixed at their respective posterior means, the OLS estimates (standard errors) of $\\delta_2$ and $\\gamma_2$ are $8.5(\\pm 0.8)$ and $-0.43(\\pm 0.08)$, respectively. The standard errors in Table~\\ref{tbl:result} are slightly larger because they account for the extra uncertainty in estimating $c_i, t_0^{(i)}, t_{\\rm max}^{(i)}$, and $t_1^{(i)}$. The quality of the predictive relationships is shown in Figure~\\ref{fig:ratio_predobs}, where we display the ratios of the values predicted from the previous cycle to those measured for that cycle, together with 1-$\\sigma$ prediction error bars.  Here the posterior mean estimates (as in Table~\\ref{tbl:fitparam}) are regarded as measured or estimated values; the predicted ones are computed using the equations in Table~\\ref{tbl:result}.  The prediction error bars are computed by simulation.  Note that the coefficients $\\hat{\\delta}_i,\\ \\hat{\\gamma}_i,\\ i=1,2,3,$ are estimated using all cycles (including those that come after the one being predicted).  \\begin{figure}[htb!] \\caption{Ratios of predicted parameter values to measured values.  The values predicted for the parameters in each cycle (amplitude, rise time and fall time) based on the parameters of the previous cycle, are compared with the values directly estimated for that cycle.} \\label{fig:ratio_predobs} \\begin{center} \\includegraphics[width=2.4in, angle=270]{cmp_pred_actual_ampl.eps} \\includegraphics[width=2.4in, angle=270]{cmp_pred_actual_rise.eps} \\includegraphics[width=2.4in, angle=270]{cmp_pred_actual_fall.eps} \\end{center} \\end{figure}  \\begin{table}[htb!] \\caption{Summary of three relationships and their parameter estimates when fitting the two stages jointly. The fitted values are posterior means and the standard errors are posterior standard deviations. } \\label{tbl:result} \\begin{center} \\begin{tabular}{lll} \\hline Cycle parameter & Relationship                                                        & Fitted value (Std.\\ err.) \\\\ \\hline Amplitude       & $c_{i+1}\\sim \\delta_1+\\gamma_1 c_i/(t_0^{(i+1)}-t_{\\rm max}^{(i)})$ &  $\\hat{\\delta}_1=4.1\\, (\\pm 1.5)$           \\\\ &                                                                     &  $\\hat{\\gamma}_1=3.9\\, (\\pm 1.0)$           \\\\ \\hline Time to maximum & $t_{\\rm max}^{(i+1)}-t_0^{(i+1)}\\sim \\delta_2+\\gamma_2 c_{i+1}$     &  $\\hat{\\delta}_2=8.5\\, (\\pm 1.0)$           \\\\ &                                                                     &  $\\hat{\\gamma}_2=-0.43\\, (\\pm 0.09)$          \\\\ \\hline Time to minimum & $t_1^{(i+1)}-t_{\\rm max}^{(i+1)}\\sim \\delta_3+\\gamma_3 c_{i+1}$     &  $\\hat{\\delta}_3= 4.3\\, (\\pm 1.5)$           \\\\ &                                                                     &  $\\hat{\\gamma}_3=0.31\\, (\\pm 0.15)$          \\\\ \\hline \\end{tabular} \\end{center} \\end{table}   \\subsection{Predictions for Cycle 24} \\label{sec:cyc24}  One advantage of a Bayesian hierarchical model in this context is that, once we obtain the samples from the posterior distribution, prediction of the characteristics of the current (incomplete) cycle is obtained automatically. Table~\\ref{tbl:cyc24} reports summaries of the posterior inference for Cycle 24, using data up to November 2008, May 2010, and October 2011, respectively.  Based on data up to October 2011, Cycle 24 is estimated to rise to maximum in January -- February 2014 $\\pm$ six months, with a maximum smoothed monthly sunspot number of $97\\pm 15$, where the estimates are posterior means and the error bars are posterior standard deviations.  (The maximum smoothed SSN, or the expected SSN at solar maximum, is $(\\beta+c_i)^2+\\sigma^2$, after accounting for the square-root transformation in Equation~(\\ref{eq:stage1}).)  It is likely to be a weak cycle with a longer-than-usual (an expected 12.1 years) total duration, although the uncertainty associated with this estimate is fairly large.  We observe that the estimated maximum smoothed sunspot number is relatively stable across the three analyses.  The large error bars associated with the November 2008 analysis highlight the inherent difficulty in making predictions before or at the onset of a cycle.  \\begin{table}[htb!] \\caption{Cycle 24 predictions based on MCMC fitting of the two stages jointly. Fitted values are posterior means and standard errors are posterior standard deviations. Max.\\ SSN refers to the smoothed monthly average sunspot number at the peak of Cycle 24. Time to rise [years] is defined as $t_{\\rm max}^{(24)}- t_0^{(24)}$. } \\label{tbl:cyc24} \\begin{center} \\begin{tabular}{rrrrrr} \\hline & $c_{24}$           & Max.\\ SSN       & Time of max.\\ [yrs]             & Time to rise          &  Cycle length        \\\\  Nov\\,08   & $9.0\\pm 1.6$       & $96\\pm 32$      &  Mar $2013 \\pm 0.98$    &  $4.7\\pm 1.1$      &  $11.9\\pm 1.7$     \\\\  May\\,10   & $8.2\\pm 1.2$       & $77\\pm 21$     &  May $2014\\pm 0.69$        &  $5.3\\pm 0.84$      &  $12.1\\pm 1.6$     \\\\  Oct\\,11   & $9.3\\pm 0.77$      & $97\\pm 15$     &  Jan/Feb $2014\\pm 0.48$    &  $4.8\\pm 0.55$      &  $12.1\\pm 1.5$     \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  Figure~\\ref{fig:cyc24} illustrates the estimates of the averages of $Y_t$ in Equation~(\\ref{eq:stage1}) (specifically, $(\\beta + U_t)^2 + \\sigma^2$). The solid curve represents the posterior mean while the upper (lower) dashed curve represents the $95\\%$ ($5\\%$) posterior quantile. Note that estimates for time points in the past are well constrained because of the available data, but future predictions are much more variable. The two-stage model is well suited for combining two pieces of information that have potential predictive power: sunspot number observations that clearly belong to Cycle 24, and the prescription of the second-stage model (Equations~(\\ref{eqn1})\\,--\\,(\\ref{eqn3})) which relates the characteristics of Cycle 24 to those of previous cycles. Given the relatively few observations at the beginning of Cycle 24, the predictions are heavily influenced by the second-stage model. As Cycle 24 progresses, direct observations will play a heavier role and the uncertainties associated with the predictions will diminish. This is illustrated by the reduction in the uncertainty band in the bottom panel which includes 35 more months of observations (up to October 2011) compared to that in the top panel (up to November 2008).  (This reduction in uncertainty is also apparent from Table~\\ref{tbl:cyc24}.) When the more recent data are included, the predictions are more driven by direct observations from Cycle 24; the fitted values are similar, but the 90\\% predictive intervals are appreciably narrower. This shows that the latest data are reasonably consistent with the second-stage relationships, and combining the two stages shrinks the error bars.  \\begin{figure}[ht] \\caption{Predictions of Cycle 24 obtained by fitting the two stages jointly, using data up to November 2008 (top), May 2010 (middle) and October 2011 (bottom).  The posterior mean of monthly average sunspot numbers is shown as the solid curve, and the 5\\% and 95\\% posterior quantiles are shown as dashed curves.  The top figure illustrates predictions for a completely new cycle, and the bottom for one that is well in progress.  May 2010 is chosen as it lies half way in between November 2008 and October 2011.  Note that the uncertainty in the predictions is reduced considerably when more data from the current cycle are included. } \\label{fig:cyc24} \\begin{center} \\includegraphics[width=2.4in, angle=270]{ssnpredict08.eps}\\\\ \\includegraphics[width=2.4in, angle=270]{ssnpredict10.eps}\\\\ \\includegraphics[width=2.4in, angle=270]{ssnpredict11.eps} \\end{center} \\end{figure}      \\end{section}  \\begin{section}{Summary} \\label{sec:summary}  We have carried out a comprehensive statistical analysis of the sunspot record. After suitably transforming the data to stabilize variance, we parameterized the shape of a cycle by its amplitude (maximum in the sunspot number), time to rise to maximum, time to fall to minimum, and the gap between its end and the start of the next cycle. By computing correlations between these parameters both within each cycle and between adjacent cycles, we have derived a set of three predictive relations.  These relations are ordered, {\\it i.e.} sequential: amplitude must be predicted first before duration and rise time. Correlations that depend on computing amplitude second are not robust, as they are subject to two influential points from early in the sunspot record (see also Vaquero and Trigo 2008).  Analyses carried out in a different order will thus lead to spurious results.  These relations can be used to predict the values of the parameters for the following cycle. We find that the best estimate for the peak in Cycle 24 is in early 2014, with an uncertainty of half a year. The maximum in the smoothed sunspot number record is expected to be $\\approx{97}\\pm{15}$, and the cycle is expected to last $12.1{\\pm}1.5$ years from the latest solar minimum (approximately November 2008). These are in the middle of the range of predictions in the literature prior to the onset of Cycle 24, and are consistent with the current estimates of the cycle parameters.  We have searched for, but do not find, any evidence for persistence beyond one cycle. There is no predictive power beyond the cycle that follows; no correlations are present, and the cycles do not retain any memory.  We also find that the cycles do not ever vanish completely. We find statistical evidence that the next cycle usually begins before the current cycle ends, as the gap between the end of a cycle and the start of a new one is usually negative.  Furthermore, we find that sunspots do not vanish entirely even if the gap were positive; however, the data are not sufficient to tell whether this holds true even in the {\\sl absence} of activity cycles."
    },
    "1208/1208.3535_arXiv.txt": {
        "abstract": "We have investigated a supra-arcade structure, associated with an M1.6 flare, which occurred on the south-east limb on 4th of November 2010. It is observed in microwaves at 17 GHz with the Nobeyama Radioheliograph (NoRH), soft X-rays in the range of 8-20 keV with the Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI), and EUV with the Atmospheric Imaging Assembly (AIA) onboard the Solar Dynamics Observatory (SDO). As reported by Reeves \\& Golub (2011), the supra-arcade structure is observed predominantly in the AIA 131 \\AA\\ channel, which contains a hot 11 MK component from Fe XIX (Boerner et al. 2011). While this hot flare plasma lasts over the decay phase of the flare, it shows some interesting characteristics in microwaves and soft X-rays: 1) In the supra-arcade structure, the brightness temperature ($T_B$) of the microwave emission increases gradually up to 2$\\times$10$^4$ K, and 2) two soft X-ray sources appear: one cospatial with the supra-arcade structure and another above the post-flare arcade. We have derived the variation of emission measure, density, and energy of the supra-arcade structure using the $T_B$ obtained from 17 GHz microwave observations. ",
        "introduction": "Recently, Reeves \\& Golub (2011) reported three events which show hot plasma above the flare arcade observed by AIA/SDO. AIA provides multi-wavelength EUV images with 1.2 $^{\\prime\\prime}$ resolution and 12 s temporal cadence. We have examined one of these events, which shows significant microwave emission in the supra-arcade structure, using NoRH, RHESSI, and AIA/SDO. ",
        "conclusions": ""
    },
    "1208/1208.0646_arXiv.txt": {
        "abstract": "Precise subtraction of foreground sources is crucial for detecting and estimating 21~cm HI signals from the Epoch of Reionization (EoR). We quantify how imperfect point source subtraction due to limitations of the measurement dataset yields structured residual signal in the dataset. We use the Cramer-Rao lower bound, as a metric for quantifying the precision with which a parameter may be measured, to estimate the residual signal in a visibility dataset due to imperfect point source subtraction. We then propagate these residuals into two metrics of interest for 21~cm EoR experiments -- the angular power spectrum and two-dimensional power spectrum -- using a combination of full analytic covariant derivation, analytic variant derivation, and covariant Monte Carlo simulations. This methodology differs from previous work in two ways: (1) it uses information theory to set the point source position error, rather than assuming a global root-mean-square error, and (2) it describes a method for propagating the errors analytically, thereby obtaining the full correlation structure of the power spectra. The methods are applied to two upcoming low-frequency instruments that are proposing to perform statistical EoR experiments: the Murchison Widefield Array (MWA) and the Precision Array for Probing the Epoch of Reionization (PAPER). In addition to the actual antenna configurations, we apply the methods to minimally-redundant and maximally-redundant configurations. We find that for peeling sources above 1~Jy, the amplitude of the residual signal, and its variance, will be smaller than the contribution from thermal noise for the observing parameters proposed for upcoming EoR experiments, and that optimal subtraction of bright point sources will not be a limiting factor for EoR parameter estimation. We then use the formalism to provide an \\textit{ab initio} analytic derivation motivating the ``wedge\" feature in the two-dimensional power spectrum, complementing previous discussion in the literature. ",
        "introduction": "An understanding of the systematic observational biases in a dataset is crucial for forming an unbiased estimate of the Epoch of Reionization (EoR) signal, as measured by low-frequency interferometric radio telescopes, because the signal is expected to be weak (orders of magnitude below contaminating signals and the thermal noise), and the scientific use of the result relies on the detailed Fourier structure of the signal. Foreground contaminants to the EoR 21~cm signal include bright point sources, confusing (unresolved) point sources, and diffuse galactic synchrotron emission. Studies have been undertaken to explore methods for removing unresolved sources and diffuse structure, and the effects of these techniques on the signal, but it is typically assumed that the bright source subtraction has been performed perfectly \\citep[][and references therein]{bowman09,liu09,liu11}. These bright sources are often used to calibrate the instrument \\citep[e.g., Murchison Widefield Array -- MWA{\\footnote[1]{http://www.mwatelescope.org}}, Precision Array for Probing the Epoch of Reionoization -- PAPER{\\footnote[2]{http://eor.berkeley.edu}}, Low Frequency Array -- LOFAR{\\footnote[3]{http://www.lofar.org}}, Long Wavelength Array -- LWA{\\footnote[4]{http://lwa.unm.edu}},][]{lonsdale09,tingay12,parsons10,stappers11,ellingson09}. Imperfect source subtraction may lead to systematic errors in the dataset used to perform EoR measurement. It is this effect we study in this paper.  There have been previous attempts to quantify the uncertainties introduced into an image dataset, and consequently a measure of the EoR power spectrum \\citep{datta09,datta10}, by position errors in the bright point source subtraction. These studies have assumed position errors with a global root-mean-square magnitude, and propagated them to the final measurement. They have not studied the expected position errors introduced by imprecise measurement of the point sources (i.e., uncertainties in parameter estimation due to imperfect data). An analysis that derives the expected parameter uncertainties, and then propagates them analytically into the visibility and image datasets, will provide the most reliable estimate of the impact of bright source contamination. In this paper we present a method for determining the impact of bright source subtraction from first principles. Our method provides a straight-forward analysis, based on the Cramer-Rao lower bounds (CRBs), of the theoretical level to which the parameters of a source, and therefore our ability to subtract it, can be determined, and propagates these uncertainties into the visibility plane and EoR power spectra.  Bright point sources affect interferometric datasets substantially. In addition to saturation of signal within a synthesized beam-width of the actual source, the incomplete $uv$ coverage of distributed synthesis instruments leads to strong sidelobe contamination throughout the dataset. For measurement of the weak ($\\sim$20--50~mK) and structured EoR signal, bright point sources ($>$1~Jy) need to be accurately and precisely removed from the dataset before analysis can proceed. Precise estimates of the source parameters are achievable due to the strong signal from these sources (high signal-to-noise ratio). Conversely, inaccurate subtraction of the brightest sources may lead to large amplitude residuals. The balance between these factors will determine the utility of these datasets for EoR science.  At low frequencies ($\\lesssim$200~MHz), the ionosphere produces a measurable perturbative effect on the wavefront shape of a propagating signal, even for short baseline arrays specifically targetting the EoR (e.g., MWA, PAPER). The ionosphere acts as a phase screen. The first order gradient term yields a shift in source position, while higher-order terms cause lensing. Spatially compact arrays see very little curvature, making a gradient phase screen a reasonable approximation under normal ionospheric conditions. \\citet{cohen09} studied the differential refraction (position change relative to other sources in the field) for Very Large Array (VLA) fields at 74~MHz. They found an empirical power law fit to the differential refraction as a function of source sky separation of $\\sim$0.5 during nighttime observations ($\\sim$0.7 during the day), under normal ionospheric conditions (refraction regime), yielding a source shift of $\\sim$100 arcseconds at a separation of 25 degrees. In addition to this, the whole field suffers an overall shift of $\\sim$1--10 arcminutes. At 150~MHz, one would expect the refraction to be half the size, owing to the frequency dependence of phase rotation.  Wide-field low frequency instruments, such as the MWA, PAPER and LWA, aim to use known bright point sources as calibrators to model the instantaneous effect of the ionosphere on the wavefront. In addition to these calibration sources, lower flux sources will need to be removed from the dataset for EoR signal estimation. In this work we consider an observational scheme where the calibration is performed in real-time, necessitating a measurement of the instrument and sky response on short timescales, and applying these measurements in real-time to the measured data (visibilities). We refer to peeled sources generically as ``calibrators'', although some sources may be peeled as contaminating foregrounds, but not used for instrument calibration.  \\citet{mitchell08} list the steps in the MWA Calibration Measurement Loop (CML) within the Real-time System (RTS). The CML is designed to track the real-time wavefront distortions due to the ionosphere, by performing instrumental calibration on short timescales ($\\sim$8--10 seconds). The calibration is performed on bright calibrators, which are then subtracted (``peeled\") from the visibility dataset before further data processing. The potential for ionospheric changes over the calibration timescale effectively corresponds to a re-measurement of the parameters of each point source each 8--10 seconds. Accurate and precise subtraction of these calibrators is therefore limited by the amount of information available about the parameters of a source (brightness and sky position) in that period, yielding a lower limit to the precision with which estimation (and therefore subtraction) can occur. For an unbiased estimator, this fundamental limit is set by the Cramer-Rao bound (CRB) on parameter estimation \\citep{kay93}. We will use the CRB to quantify the maximal precision with which the parameters (sky position and flux density) of a source may be estimated, and use these uncertainties to model the residual point source signal in the visibility dataset. We then describe a formalism for propagating these errors to power spectra, using a fully-covariant analytic derivation. We use this formalism, along with Monte-Carlo simulations to estimate the residual power and uncertainty in the angular power spectrum and two-dimensional power spectrum due to point source subtraction. Note that this approach differs from that of \\citet{datta10} in two ways: (1) it uses information theory to set the point source position error, rather than assuming a global error; (2) it propagates the full covariance matrix to the power spectra estimates, yielding both the uncertainties on the power values, and the level of correlation between $k$-modes. We work from first principles in our approach. ",
        "conclusions": "Our results indicate that (1) optimal point source subtraction to 1~Jy will not be a limiting factor in statistical EoR estimation for proposed experiments with MWA and PAPER, contrary to the conclusions of previous analyses \\citep{datta10}, which relied on assumed position errors that were independent of source strength; (2) frequency-channel covariance combined with spread of visibilities across coherent $(uv)$ cells yields the ``wedge\" feature observed in point source subtracted EoR simulations; (3) angular mode correlations are expected to be small (1--2 orders of magnitude below variance level). We have presented an end-to-end derivation of the approximate shape of the wedge feature, using the first principles framework we have developed. Here we describe some interesting features and caveats of these results, based on the assumptions within our framework.  The magnitude of the residual point source signal is dependent on the number of calibrators being peeled, the thermal noise level using the full array and bandwidth, and the angular mode in question (with a ${\\sim}|\\boldsymbol{u}|^2$-dependence on variance in the visibility). The uncertainty in the power spectrum at angular mode $l$ therefore scales with: \\begin{equation} \\sigma_{\\rm PS}(l) \\sim \\frac{N_{\\rm cal}l\\sigma_{\\rm therm}\\Delta\\nu}{\\Delta\\nu_{\\rm tot}} \\propto \\frac{N_{\\rm cal}l}{N_{\\rm vis}\\Delta{t}\\Delta\\nu_{\\rm tot}}, \\end{equation} where $\\sigma_{\\rm therm}$ is the thermal noise level across the entire bandwidth, $\\Delta\\nu_{\\rm tot}$, for power spectrum channels of width $\\Delta\\nu$. The magnitude of the expected power, and power uncertainty, therefore scale linearly with the number of peeled sources, and inversely with the total system bandwidth and experiment time. The PAPER array suffers from 1/4 the visibilities of the MWA and a larger field-of-view, but this is balanced somewhat by the additional bandwidth and the smaller $l$ modes available to the (more compact) instrument.  We have assumed a sequential peeling scheme, whereby sources are independently and sequentially estimated, and peeled from the dataset, without downstream impact on the estimation of the parameters of other sources. In reality, an experiment can choose to either (1) estimate the parameters of all sources simultaneously, whereby the approximation of pure radiometric noise in the dataset (ignoring calibration errors) is appropriate, or (2) estimate and peel sources sequentially, whereby the covariant noise matrix at each iteration is a combination of uncorrelated radiometric noise and correlated residual signal noise. The choice between these estimation techniques depends on the number of sources being peeled relative to the available data (number of unknowns versus constraints; the CRB is singular for underconstrained problems), and we leave this discussion for future work. If we ignore correlations between visibilities, a simple back-of-the-envelope calculation demonstrates that, for the parameters used in this work, our assumption is reasonable (where we have effectively expanded the noise covariance matrix as a Taylor series, and kept the first linear term). For subtraction of a larger number of sources, or for fewer visibilities, this approximation breaks down, and the compound effect of imprecisely estimating source parameters within ever-noisier data will increase the residual signal power towards the amplitude of the thermal noise power.  A related assumption in our model is the ignorance of previous calibration solutions for estimation of the parameters of a source. As demonstrated in Figure \\ref{source_precision_plot}, optimal source position estimates with the full-bandwidth PAPER array (and an 8 second integration) have precision comparable to the potential ionospheric differential refraction magnitude, for sources close to the 1~Jy peeling floor. In these cases, the utility of the current dataset may be augmented by use of prior information about the local behaviour of the ionosphere. For larger arrays, this is also an issue if one attempts to peel below $\\sim$1~Jy. The optimal method for estimating source location (balance of prior and current information, and how these are combined) will depend on the peeling floor, the efficiency of the estimation method, and the specific array design.  Discussion of realized estimation performance highlights a further basic assumption of our framework; we have proceeded from a fundamental reliance on the existence of an efficient and unbiased estimator (i.e., an estimation method that achieves the CRB on parameter precision). This has allowed us to calculate \\textit{lower limits} on the impact of point source subtraction, and using a first principles approach, but does not address the issue of the existence of such a method (a method may not exist that can achieve the bound). In addition to designing an efficient estimator, the CRB applies only to unbiased estimators, and introducing bias into the estimation will yield different structures in the final power spectrum than shown here. As a simple estimate of the impact of an unbiased, but inefficient estimator, we (arbitrarily) assume we have estimated the source parameters to within three times the CRB, and calculate the residual signal uncertainty in the angular power spectrum. This toy model yields an increase in uncertainty of a factor of $\\sim$10, indicating a rough square scaling of estimator inefficiency with noise uncertainty.  We have derived a ${\\sim}|\\boldsymbol{u}|^2$-dependence of visibility-based variance on baseline length, demonstrating a ramping of noise power at large angular modes (small angular scales). In this work, we have used the same antennas to estimate the source parameters and produce the power spectra (the only difference being the bandwidths used). The long baselines contribute relatively more than the short baselines to the estimate of source position (equations \\ref{l_precision}--\\ref{m_precision}). However, these baselines form the largest $l$ modes, where residual signal has the greatest impact. Inclusion of these baselines for parameter estimation is crucial for precise estimates; their use in EoR measurements depends on the experimental design, and the desire for measuring the 21~cm EoR signal at those scales. Use of these baselines does not affect EoR estimation at other scales (the correlation between modes is small). Also, as demonstrated in \\citet{trott11}, inclusion of short baselines for position estimation is important (but less so than long baselines), due primarily to the increased instrumental sensitivity when using all available antennas.  Our results also demonstrate the expected behaviour of minimally-redundant versus maximally-redundant arrays. Uniform sampling yields even thermal noise power across a wide range of angular modes, while maximally-redundant arrays sample some scales well, but yield patchy performance and restricted range of sampled modes. Choice between these models depends on the experiment being performed, but a maximally-redundant array requires prior information about the most interesting scales for estimating the 21~cm signal, while the uniform arrays are a better ``starter\" design, where the scientist has no prior information about signal structure and is aiming for a first detection/estimation."
    },
    "1208/1208.5313_arXiv.txt": {
        "abstract": "This short note points out two of the incongruences that I find in the \\citet{loredo12} comments on \\citet{andreon2012bayesian}, i.e. on my chapter written for the book  ``Astrostatistical Challenges for the New Astronomy\". First, I find illogic the Loredo decision of putting my chapter among those presenting simple models, because one of the models illustrated in my chapter is qualified by him as ``impressing for his complexity\". Second, Loredo criticizes my chapter at one location confusing it with another paper by another author, because my chapter do not touch the subject mentioned by \\citet{loredo12} critics, the comparison between Bayesian and frequentist fitting models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.3253_arXiv.txt": {
        "abstract": "In this paper we present a spectral atlas covering the wavelength interval 930--1188\\,\\an\\ for O2--O9.5 stars using {\\it Far Ultraviolet Spectroscopic Explorer} archival data.  The stars selected for the atlas were drawn from three populations: Galactic main sequence (class III-V) stars, supergiants, and main sequence stars in the Magellanic Clouds, which have low metallicities. For each of these stars we have prepared FITS files comprised of pairs of merged spectra for user access via the {\\it Multi-Mission Archives at Space Telescope.} We chose spectra from the first population with spectral types O4, O5, O6, O7, O8, and O9.5 and used them to compile tables and figures with identifications of all possible atmospheric and ISM lines in the region 949-1188\\,\\an.~ Our identified line totals for these six representative spectra are 821 (500), 992 (663), 1077 (749), 1178 (847), 1359 (1001), and 1798 (1392) lines, respectively, where the numbers in parentheses are the totals of lines formed in the atmospheres, according to spectral synthesis models. The total number of unique atmospheric identifications for the six main sequence O star template spectra is 1792, whereas the number of atmospheric lines in common to these spectra is 300. The number of identified lines decreases toward earlier types (increasing effective temperature), while the percentages of ``missed\" features (lines not predicted from our spectral syntheses) drops from a high of 8\\% at type B0.2, from our recently published B star far-UV atlas, to 1--3\\% for type O spectra. The percentages of overpredicted lines are similar, despite their being much higher for B star spectra. We discuss the statistics of line populations among the various elemental ionization states. Also, as an aid to users we list those isolated lines that can be used to determine stellar temperatures and the presence of possible chemical anomalies. Finally, we have prepared FITS files giving pairs of merged spectra for stars in our population sequences for access via the {\\it Multi-Mission Archives at Space Telescope.} ",
        "introduction": "\\label{int}  High dispersion spectroscopy provides windows into the past and current physical processes in massive O stars as clear as for stars anywhere on the H-R Diagram.  Individual O stars can be found in different stages of evolution because of their short lifetimes and unique spectral signatures that advertise these stages. Indeed, many of them evolve to SN\\,Ib and SN\\,Ic supernovae by first becoming Luminous Blue Variables and/or Wolf-Rayet stars, which are easily identifiable by their light curves or broad spectral emission features. Likewise, O stars in close binaries are readily discovered from emission in optical and UV lines excited by interacting winds. Because the incidence of lines in O star spectra and continuum flux both peak in the far-ultraviolet wavelength region, the use of space-borne instrumentation is critical to fully understanding their atmospheres and radiative processes. Although observations with the {\\it Far Ultraviolet Spectroscopic Explorer} satellite (\\fuse) were terminated in 2007, the \\fuse\\ archive now offers a large sample of stellar spectra having identical wavelength coverages provided by homogeneous data processing. In particular, \\fuse-accessible O stars reside not only within the solar neighborhood in the Galaxy but also in the more distant Magellanic Clouds. Because their lifetimes are brief, O stars provide an invaluable record of the recent chemical composition and angular momentum histories in the solar neighborhood and the return of much of their nuclear-processed matter to the Interstellar Medium (ISM).  Astronomers have historically been quick to capitalize on uniformly reprocessed data held in archives following the close of UV spectroscopic mission.  The availability of these archives has allowed the construction of a number of fine spectral atlases recorded by the {\\it Copernicus} and the {\\it International Ultraviolet Explorer ( IUE}) satellites. Important examples of these compilations for the middle-UV spectral range are the {\\it Copernicus} atlas of $\\tau$\\,Sco (Rogerson et al. 1978) and {\\it IUE} pictorial atlases of O and B stars (Walborn, Nichols-Bohlin, \\& Panek 1985, Rountree \\& Sonneborn 1993, Walborn, Parker, \\& Nichols 1995). Pellerin et al. (2002) inaugurated a second group of digitized far-UV (\\fuse) atlases consisting of representative spectra of O and early-B type, luminosity class I-V stars in the Galaxy and Magellanic Clouds. These atlases have been supplemented by figures constructed by N. Walborn for the Gray \\& Corbally (2009) monograph on stellar spectral classification that show the changes of primary spectral features in the far-/middle-UV with luminosity class in main sequence O stars.  Such atlases have been effective in showcasing the general trends of the strong photospheric lines and the so-called ``UV resonance wind\" lines with effective temperature T$_{\\rm eff}$ and luminosity class. For example, a reconnaissance of the O-star (Pellerin et. al) atlas shows how far-UV spectra of most O stars are strongly mutilated by interstellar lines, particularly below 1100\\,\\an. The photospheric component of these spectra is dominated by the confluence of high series Lyman lines, Ly$\\beta$--Ly$\\theta$ down to the blue limit of the \\fuse\\ instrument at 920\\,\\an.~ The strongest metallic feature formed substantially in the photosphere is generally the C\\,III line complex at 1174--1176\\,\\an. This multiplet arises from an excitation of the lower atomic levels at an excitation  $\\chi$$_{exc}$ = 9\\,eV.~ In supergiant O stars radiative winds can cause this complex to develop into a broad P\\,Cygni profile. Other strong features are resonance lines of C\\,III (977\\,\\an), N\\,III (991\\,\\an), O\\,VI (1031\\,\\an, ~1037\\,\\an), Si\\,IV (1128\\,\\an), S\\,IV (1062\\,\\an,~ 1072\\,\\an,~ 1073\\,\\an,~ 1098--1100\\,\\an), P\\,IV (950\\,\\an), and P\\,V (1117\\,\\an,~ 1128\\,\\an). Spectral lines of O stars suffer doppler broadening, and this serves to wash out isolated weak metallic lines that would otherwise be resolved in high dispersion spectra. However, as detailed below, the far-UV spectra have the saving grace that, except for the liberally distributed ISM lines, the photospheric lines tend to be more widely spaced and to blend less with one another than in the middle-UV wavelengths.  The Pellerin et al. atlas and Barnstedt et al. (2000) have displayed valuable detailed information about the presence of molecular H$_2$ features formed in the ISM. From the point of view of stellar atmosphere investigators, these features contaminate many of the far-UV lines formed in the atmospheres of most Galactic stars. The Pellerin et al. atlas was closely followed by a second spectral atlas of OB stars in the two Magellanic Clouds (Walborn et al. 2002).  This work concentrated on the behavior of wind lines with respect to metallicity as well as effective temperature and luminosity.  In addition, Blair et al. (2009) published a compendium of {\\it FUSE} spectra of hot stars in the Magellanic Clouds. This atlas focused on the identification of far-UV resonance lines formed in the ISM, including those found at multiple velocities.  The first far-UV high-dispersion spectral coverage published for a B star was the {\\it Copernicus} atlas of $\\tau$\\,Scorpii (B0.2\\,V; Walborn 1971) by Rogerson \\& Upson (1978). Rogerson \\& Ewell (1985; ``RE\") published a detailed tabulation of atmospheric and ISM lines identified in this atlas. The RE work was undertaken at a time when spectral synthesis tools were not commonly available, and when only a relative handful of experienced spectroscopists who were also specialists in atomic physics could make reliable line identifications. Even so, in the absence of commonly available synthesis programs at that time, it was difficult to make wholesale line identifications without some errors. This situation has changed dramatically in the intervening years with the development of spectral line synthesis tools that make use of extensive atomic line libraries.  Inspired by the {\\it Copernicus} atlas, Smith (2010; hereafter ``Paper\\,1\") constructed a far-UV spectral atlas for B stars using spectra from \\fuse\\, {\\it IUE,} and the {\\it Space Telescope Imaging Spectrograph} ({\\it STIS}) over the (vacuum) wavelength range 930--1225\\,\\AA.~ In making this atlas we set an arbitrary short wavelength limit of 930\\,\\an\\ because, other than for a few blended Lyman lines, no useful information about the photospheric spectrum could be obtained below this wavelength. We chose the red limit by including spectra from the {\\it HST/STIS} and/or {\\it IUE} cameras in order to cover the Lyman\\,$\\alpha$ feature and other important lines in its vicinity.  The B star atlas addressed its first goal of fulfilling the need for the identification of all possible visible lines in the ``template\" spectra of three representative main sequence B0.2\\,V, B2\\,V, and B8\\,V stars. These identifications were made by using published atomic line libraries that predict occurrences of lines from spectral synthesis models. This atlas also recorded ``misses,\" i.e., the wavelengths of observed lines that could not be predicted from our line syntheses.  The rotational velocities of stars selected for the atlas are relatively low in order to resolve as many neighboring photospheric lines as possible. Therefore, we referred to this work as a ``detailed spectral atlas.\" The second goal of this atlas was to provide uniform spectral data products to the astronomical community through the wavelength range just described.  The numbers of lines identified for the three template spectra in Paper\\,1 were 2288 (2004), 1612 (1465), and 2469 (2260) lines, respectively. Here the values given in parentheses are the number of lines found within almost the full wavelength range covered by the \\fuse\\, 949--1188\\,\\an.~ Of these totals only small percentages (8\\%, 2\\%, and 2\\%, respectively) of photospheric lines could not be identified, which is to say that the oscillator strengths (log\\,$gf$'s) for these lines are at best poorly determined and are thus not included in our line library. The B star atlas results were published more fully in the electronic edition, The spectral data files and all other products were made available for public download as one of MAST's\\footnote{The Multimission Archives for Space Telescopes is located at the Space Telescope Science Institute (STScI). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract.  Support for archiving MAST data is provided by NASA Office of Space Science under grant NAS5-7584.} ``High Level Science Product\" (HLSP) area (http://archive.stsci.edu/prepds/fuvbstars/).  This paper presents a similar spectral atlas for O stars. The atlas is organized in much the same way as the B star atlas. However, one important difference is that whereas the blue wavelength limit we chose, 930\\an, is the same as for the B star atlas, the red one is set by the \\fuse\\ spectral coverage, again, $\\approx$1188\\an. In particular, we found for the present atlas that it was not easy to again include the 1188--1225\\an\\ region that was previously surveyed in the B star atlas because of the paucity of O stars observed systematically in this spectral region.  Our presentation is organized as follows. The selection of spectra and the methodology for data handling, including the creation of FITS spectra for all stars in our sample, as well as for the line identification in six exemplars, are described in $\\S$2. In $\\S$3 we display portions of the atlas and give a detailed list of several thousand identifications as well as a list of ``clean\" lines across much of the O-star domain. In $\\S$4 we give relevant statistics from our identifications for the six O4\\,V, O5\\,V, O6.5\\,V, O7\\,III, O8\\,V, and O9.5\\,III exemplars we call our O star spectral templates. We also comment in detail on a number of possible spectral markers for physical conditions (mainly effective temperatures) in these stars' atmospheres, all of which are assumed to be in hydrostatic equilibrium. ",
        "conclusions": "\\subsection{Line statistics} Even a glance at the Tables 3 and 4 suggests that the density of photospheric lines, e.g., those in the range 949--960.0\\an, decreases towards earlier types. For the same wavelength range the number of lines for the B0.2 template was 81 (Paper\\,1), and in our tables it is 53 for O9.5\\,V and 24 for O4\\,V. Near the long wavelength limit, 1170-1188.5\\,\\an, the corresponding decline in numbers of identifications is as least as dramatic, from B0.2 to O9.5 to O4: 128, 97, and 31, respectively. Another important generalization is that with few exceptions the strongest features in the observed spectra are ISM or hydrogen Lyman lines -- that is, nearly all metallic lines are weak. One can get a superficial impression that strong photospheric features appear in far-UV spectra of earlier type stars, but this is due to saturation of wings of strong H$_2$ features because of longer ISM column lengths to these distant stars. The numbers of all identified lines in our Tables\\,3 and 4, including ISM ones, total 1798, 1359, 1178, 1077, 992, and 821, for O9.5, O8, O7, O6, O5, and O4, respectively. When we separate out the ISM lines the totals for photospheric lines alone are 1392, 1001, 847, 749, 663, and 500, respectively. The totals are depicted by a thick dotted line in Figure\\,\\ref{frc}. As one proceeds to earlier O spectral types, the downward trend shown by these totals contrasts with the nonmonotonic numbers of lines found along the B spectral sequence in Paper 1.  In the comparable far-UV wavelength region those numbers were 2182 (B8), 1398 (B2), and 1991 (B0.2). Not surprisingly, one finds that there is a fair degree of overlap among lines of neighboring spectral subtypes. The total number of unique photospheric identifications is 1792. The number of photospheric lines we identified for all six O spectral types is 300. Note that because our stellar line identifications are based on line synthesis predictions (except for missed features seen in the spectra), our O star statistics are not affected by the increased line broadening in our template spectra with respect to the (generally smaller) line broadening encountered for our B star spectra presented in Paper\\,1.  \\begin{figure} \\centering \\includegraphics[,scale=.35,angle=90]{f3.eps} \\vspace*{0.2in} \\caption{ The run with spectral type of the fraction of identified far-UV photospheric lines for O stars near the main sequence. Fractions of lines for various elements are differentiated. Nearly 80\\% of all lines arise from C, N, O, Cr, or Fe ions. Line and symbol types are used just to clarify these relations. The dotted line is the fraction of photospheric lines identified for each spectral type relative to the 1392 lines in the O9.5 spectrum. Numbers under the O9.5-type fraction are the maximum {\\it number} of lines for the element(s) for all O spectral types.} \\label{frc} \\end{figure}   The fundamental cause of the changes in these line totals is the shifting ionization states with increased effective temperature. This is easiest to see for lines of the iron group elements because their differences in ionization potential are relatively modest for successive ion states. As we noted in Paper\\,1, the changes in B star spectra occur because the total number of visible Fe\\,II plus Fe\\,III lines in the B2 far-UV spectrum is lower than the totals of Fe\\,II plus Fe\\,III lines for either B0.5 or B8 spectra. Thus, early and late type B spectra exhibit more lines that are mainly either Fe\\,II lines or Fe\\,III lines, respectively, than the sum of the roughly equal numbers of Fe\\,II and Fe\\,III lines in a B2 spectrum. This type of ``rule\" also holds true for the rapidly changing Fe ionization states among types O9.5 through O6. In fact, the O9.5 spectrum exhibits a total of 656 Fe lines, about half (324) of which are Fe\\,III lines. However, as we progress to O8, O7, and O6 types, Fe$^{4+}$ has become the dominant state and 60-65\\% of the visible iron lines are Fe\\,V. By contrast, when we reach subtypes O5 and O4, even though the dominant ion stage has shifted from Fe$^{4+}$ to Fe$^{5+}$, most visible iron lines in the far-UV are still Fe\\,V lines. The ultimate cause of these changes in our statistics in the far-UV is that the visible Fe\\,VI and Fe\\,VII lines have moved into the extreme UV as the typical energy differences between transition energy levels increases with ionization potential. Extreme-UV lines are no longer visible to us, owing to the high ISM opacity at Lyman continuum wavelengths. Also, the dominant carbon and nitrogen (and to a lesser extent oxygen) atoms have been stripped of all or nearly all of their electrons and contribute relatively few lines at far-UV wavelengths.  Another attribute of our line statistics is the similar percentages of secondary photospheric lines relative to the total membership of a line group remains about the same:~ 61\\% for O8--O9.5 and 64-66\\% for O4--O6. In B stars (Paper\\,1) the percentages are roughly the same, 59--66\\%.\\footnote{We omit the $\\tau$\\,Sco identifications from this statement because in Paper\\,1 we used a narrower wavelength window, ${\\pm 0.05}$\\,\\an,,~ for assignment of lines to co-blended line ``groups.\"} This means that the degree of blending of photospheric lines within predominantly photospheric line ``groups\" is similar for both B and O spectra in the far-UV.   We were pleased to discover that the percentage of unknown (i.e., ``UN~I,\" or underpredicted) lines found relative to successfully predicted photospheric lines decreased in O star spectra relative to the B star spectra in Paper\\,I. These percentages come to 1\\% for O7--O9.5 stars, jumping only to 3\\% for O4--O6. On the other hand, the fractions of overpredicted (predicted but not observed) lines in the B star lines were high (6--12\\%). We find that this population declines from 3\\% at O9.5 to a minimum of near 1\\% for O6--O8 stars and increases again to more than 3\\% for O4--O5. These low percentages are all the more surprising when one recalls that the ``overpredicteds\" are features we have {\\it observed} in the spectrum, whereas our ``identified\" features are {\\it predicted} from synthesis models -- whether they are visible in spectra or not. These percentages also speak well to the relative completeness of the published libaries for far-UV lines arising from multiply ionized ion states.  \\begin{figure}[ht!] \\centerline{ \\includegraphics[scale=.35,angle=90]{f4.eps}} \\vspace*{0.2in} \\caption{ The run of Fe\\,III through Fe\\,VI lines with spectral type for identified photospheric lines in the B star and O star (this paper) atlases. For ease of reference, each line is scaled to the same maximum as the Fe\\,V population. The numbers following the * symbol are these scaling factors. } \\label{fenums} \\end{figure}  The general trend of line populations arising from the most prolific line-producing elements in the far-UV spectrum, namely Fe, Cr, C, N, and O, is exhibited in Fig.\\,\\ref{frc}. Notice from this plot that 75--80\\% photospheric lines in the far-UV are caused by these five elements. The fractional contributions from these elements to the total number of lines remains nearly the same with spectral type. The exceptions are the numbers of N and O features: the numbers of N\\,III and O\\,III lines drop off at O5--O4 because of shifts in ionization. By comparison, the numbers of N\\,IV and O\\,IV--O\\,VI lines in spectra of these hotter stars are negligible.   Fig.\\,\\ref{frc} also shows that iron ions alone contribute at least half the visible photospheric lines in far-UV spectra of O stars. This is somewhat higher than the corresponding fractions (34\\%--47\\%) in the B star atlas.  Table\\,5 of that atlas shows the relative numbers of lines from various iron ions. As already mentioned, we found that the numbers of Fe\\,III (and to a lesser extent Fe\\,IV) lines peak at B0.2. Similarly, the Fe\\,II lines were found to peak at B8. In Figure\\,\\ref{fenums} we present combined totals for all iron ions that contribute to the far-UV spectra of B stars and O stars.  This plot shows that whereas the ``big story\" for B star spectra was the transition from Fe\\,II to Fe\\,III lines, for O star spectra it is the dominance of Fe\\,V lines. This occurs once the transition from Fe\\,III to Fe\\,V lines sets in,  at subtype O9.5. Fe\\,II lines are nonexistent in the photospheric spectra of O stars.    \\subsection{Temperature and chemical Indicators} \\label{chem}  We end our description of this atlas by discussing specific far-UV lines that can be used in O star main sequence spectra to define new diagnostics of effective temperature and in some instances chemical composition. For O stars many of the most interesting abundance anomalies arise from interior CNO-cycle processes. However, iron group abundance differences could also be affected by reprocessing in recent supernova explosions. Therefore, we also explore the isolated lines already introduced in Tables\\,5 and 6 for a number of ions. For the star codes of these tables, decimal values ending in ``.1,\" such as ``1.1,\" signify that a dominant line is closely blended with a line of the same ion, i.e., the blend is effectively one strong line. We refer to these as ``self-blended\" features.   \\noindent {\\it Diagnostic lines:}  \\noindent {\\bf He\\,I:~}  All five He\\,I lines in Table 5 are visible in all O template spectra, but most of them are severely blended with local and often broad H$_2$ features. The best T$_{\\rm eff}$ indicators are the lines at 958.7\\,\\an\\ and 1084.9\\,\\an. The strengths of both peak at O6--O7. \\\\   \\noindent {\\bf C\\,II:~}  We found only one completely isolated C\\,II line, 1010.371\\,\\an,~ in our spectra of O9.5--O6\\,V dwarfs. A weaker member of this 6\\,eV multiplet at 1010.083\\,\\an\\ can sometimes be discerned in the red wing of a strong H$_2$ feature at 1010\\,\\an\\ in redshifted spectra. Two other features at 1065\\,\\an\\ are present in narrow-lined spectra of late O stars, but they are overwhelmed by local H$_2$ features. We do not recommend attempting to use C\\,II lines for quantitative analyses of O star spectra.  \\noindent {\\bf C\\,III:~}  The C\\,III lines peak at subtype O9.5 in our main sequence template spectra. A single line at 1070.331\\,\\an\\ is optically thin but lies within 0.3\\,\\an\\ of a H$_2$ feature. Therefore it cannot be used for studies of carbon abundances, at least for stars having small Doppler shifts. With the exception of 977.020\\,\\an,~ the C\\,III lines shown in Table\\,5 (1125\\,\\an--1176\\,\\an) are blended with other weak multiplet members. As we pointed out in Paper\\,1, the wings of the 977\\,\\an\\ line are sensitive to electron pressure for all O spectral subtypes and thus can be used with the C\\,III 1175.9--1176.3\\,\\an\\ multiplet lines to determine a star's luminosity class.  According to Table\\,5 and Figure\\,\\ref{plt1}, the wings of the  1174--1176\\,\\an\\ complex can be blended by Fe\\,IV and Fe\\,V lines at types O7--O9.5. Thus, equivalent width ratios measured from these lines and 977\\,\\an\\ are best formed by including only the cores of the broad 1174-1176\\,\\an\\ features. Finally, as already noted, in supergiant spectra the C\\,III complex has transitioned to a P\\,Cygni profile, and the photosphere no longer makes a significant contribution.  \\noindent {\\bf C\\,IV:~} The most visible lines of C\\,IV lie close together at 1168.847\\,\\an\\ and 1168.999\\,\\an\\ and so may often be treated  as a single feature. Thus, the strength of the feature peaks at subtypes O6--O7. In Paper\\,1 we noted that the 1168\\,\\an\\ line makes its appearance at subtype B0.2, and along with the nearby C\\,III complex can form a good diagnostic if the temperature is isolated by lines of other ions. This is also true in the late O stars for classes III-V. As one proceeds to hotter temperatures the C\\,IV line becomes considerably stronger, making it more sensitive to temperature than luminosity. A much weaker isolated feature lies at 1097.319\\,\\an\\ and begins to fade at O5.  \\noindent {\\bf N\\,III}: Two isolated N\\,III lines are visible at 1103.044\\,\\an\\ and 1106.036\\,\\an, but they are very weak and therefore are not good diagnostics.  Both 1112.648\\,\\an\\ and 1152.406\\,\\an\\ lines lie close to ISM features and therefore can be used at best with caution. The best N\\,III diagnostics for temperature or abundance in the far-UV are 1183.032\\,\\an\\ and 1184.514\\,\\an.~ However, they are self-blended by other N\\,III lines with similar log\\,$gf$ values.  \\noindent {\\bf N\\,IV:~} Two isolated N\\,IV lines are visible at 1131.488\\,\\an\\ and 1133.121\\,\\an.~ An advantage provided by these lines is that their strengths do not vary significantly with spectral type. Therefore the lines can serve as nitrogren abundance indicators for all luminosity classes.  \\noindent {\\bf O\\,III:~} Lines of this ion are numerous in the optical-UV spectra of B stars, but they nearly disappear in O star spectra. Nonetheless, the oxygen and iron group abundances in these stars tend to track well, and a few remaining O\\,III lines can be used, especially with Fe\\,V lines, for temperature determinations. The combined strengths of three isolated O\\,III lines at 1007--1008\\,\\an\\ in practice renders them a single feature even in moderately broadened (v$sin\\,i$ $\\sim$ 100\\,km\\,s$^{-1}$) O type spectra. The best lines for diagnostics are those that are self-blended, such as 1149.602\\,\\an\\, 1150.882\\,\\an\\, 1153.022\\,\\an\\, 1160.154\\,\\an\\, and 1172.451\\,\\an,~ and two isolated lines at 1157.041\\,\\an\\ and 1157.161\\,\\an.~ However, even the latter lines generally overlap due to the velocity broadening.  \\noindent {\\bf O\\,IV:~} Although less sensitive to temperature than O\\,III lines, the excitation potentials $\\chi$$_{exc}$ of O\\,IV lines are in the range 48--63\\,eV. The lines increase in strength as one goes to earlier types.  The lines at 1045.364\\,\\an,~ 1046.313\\,\\an\\,~ and 1080.969\\,\\an\\ are isolated but weak. The lines 1067.832\\,\\an,~ (self-blended) and 1067.959\\,\\an\\ are otherwise isolated, but they may blend from velocity broadening. The 1080.969\\,\\an,~ 1164.546\\,\\an,~ and 1167.478\\,\\an\\ lines are good diagnostics too, but they arise from higher excitations and disappear at subtype O4.  In terms of strength and isolation, the best features among these candidates are the two 1067--1068\\,\\an\\ lines.  \\noindent {\\bf O\\,V:~} Three O\\,V lines given in Table\\,5, 995.087\\,\\an,~ 1010.602\\,\\an,~ and ,1090.320\\,\\an\\ all appear weak in our O main sequence spectra, but otherwise they are viable temperature diagnostics for most of our template spectra. Because the ionization fraction decreases with advanced spectral type, the O\\,V lines become relatively weak at O9.5\\,V. At this subtype 995.087\\,\\an\\ becomes dominated by neighboring Fe\\,III lines and cannot be used. We remind the reader that the typical $\\chi$$_{exc}$'s for most O\\,III lines in our sample lie in the range 24--36\\,eV. The $\\chi$$_{exc}$ is single-valued at 92\\,eV for O\\,V lines. As long as O\\,III and O\\,V lines can be found in an O star spectrum, their ratios can provide good measures of the ionization in the photosphere.  \\noindent {\\bf O\\,VI:~} The 1031.926\\,\\an\\ and 1037.617\\,\\an\\ resonance doublet lines have been known as critical diagnostics of mass loss in O stars and B supergiants since their discovery from UV rocket experiments long ago. Their anomalous strengths are due to the ``superionization\" of this ion, a phenomemon which has long been ascribed to the atomic Auger effect from X-ray irradiation. However, it now appears instead that the strengths may be due to clumping of parcels within a tenuous component of the O star winds (Zsarg\\'o et al. 2008). The doublet is badly marred by local H$_2$ and nearby Fe\\,V lines, and for this reason we do not list them in Table\\,5 as isolated lines for main sequence stars. Their strengths grow rapidly with temperature and luminosity, as O$^{5+}$ becomes the dominant ionization state. Save for late O main sequence spectra, the wind contribution dominates the formation of the resonance lines across most of the profiles. Their wings are typically in emission and asymmetric.  \\noindent {\\bf Si\\,III:~} Visible lines of both Si$^{2+}$ and Si$^{3+}$ have the advantage of being strong although they are nearly all self-blended. For practical purposes the sole representative Si\\,III line is 1108.358\\,\\an.~  \\noindent {\\bf Si\\,IV:~} Except for minor contaminations by local weak Fe features, the Si\\,IV 1066.614\\,\\an\\,, 1122.485\\,\\an,~ and 1154.621\\,\\an\\ lines can be used along with Si\\,III as temperature diagnostics. Since the $\\chi$$_{exc}$'s for these lines are $\\le$31\\,eV, these lines form a low-excitation complement to the more excited oxygen lines as diagnostics of photospheric temperature.  \\noindent {\\bf P\\,IV:~} To varying degrees all the P\\,IV lines in Table\\,5 are isolated in our O spectra, even given moderate velocity broadening. The strengths of all these lines grow with P$^{3+}$ ionization fraction and atmospheric temperature. With two caveats the best lines for temperature diagnostics in main sequence O star spectra are 1118.552\\,\\an\\ and 1187.540\\,\\an.~ The first caveat is that the 1118\\,\\an\\ line is one of the lesser isolated lines in this small group and is often overwhelmed by a strong P\\,V neighbor. The second caveat is that the instrumental sensitivity of the single LiF Side\\,1B {\\it FUSE} detector is low where the 1187.540\\,\\an\\ line is recorded. We point out that the 1187\\,\\an\\ line can often be seen in well exposed {\\it IUE} and {\\it HST} spectra of O stars. Also, from our line synthesis we believe the line displayed at the position of 1187.5\\,\\an\\ in the Brandt et al. (1998) atlas of the O9\\,V star 10\\,Lac is actually this P\\,IV resonance line and not an Fe\\,V line they identified.  We also note that the atmospheric contribution to the P\\,IV resonance line at 950.657\\,\\an\\ is stronger than the ISM component, at least in most of our template spectra. In our spectral atlas identifications for O7--O9.5 stars the ISM component of this same transition is blueshifted to an apparent position of $\\approx$1150.45\\,\\an.~ We can see there that the ISM component is weaker than the same line in the photosphere of the O8 template spectrum (see continuation of Figure 1 in the on-line edition). This is very likely true for the O9.5 and O7 spectra as well.  \\noindent {\\bf P\\,V:~} As we know from the Pellerin et al. (2002) O star atlas, the P\\,IV 1118.043\\,\\an\\ and 1128.010\\,\\an\\ are often among the strongest features in the far-UV spectra of O stars. The wind contribution is especially important. For example, in their detailed fittings of the profiles of these lines, Fullerton, Massa, \\& Prinja (2006) noted that near O7 the photospheric component even near the line core is generally dominated by the wind component except in those cases where the wind is weak.  This fortunate circumstance occurs because of the low abundance of this low-alpha element. Substantial P\\,Cygni emission can be seen in the wings of supergiant spectra among all subtypes O4--O9.5. Both resonance lines are strong, and yet they are usually optically thin in the winds of luminous O stars. These conditions make them superb diagnostics of mass loss from radiative winds. Fullerton et al. also found that a nearby Si\\,IV 1128 line contributed to the red wing of this doublet substantially in O7 spectra for all luminosity classes.  It should be noted that a weak P\\,V line at 1000.358\\,\\an\\ is also weakly present in all our template spectra.  \\noindent {\\bf S\\,IV:~} The most important S\\,IV lines are the resonance lines at 1062.662\\,\\an,~ 1072.974\\,\\an,~ and 1073.516, which are exhibited in the Pellerin et al. atlas. All of these attain maximum strengths at about O7.  Pellerin et al. also noted the presence of S\\,IV lines at 1098.362\\,\\an,~ 1098.929\\,\\an\\,, 1099.482\\,\\an,~ and 1100.051\\,\\an.~ These lines are visible up to O6 in our main sequence templates, even though nearby H$_2$ contributions often overwhelm the 1099\\,\\an\\ and 1100\\,\\an\\ lines. ``Also rans\" as temperature diagnostics are the weak 1106.487\\,\\an\\ line and the  stronger 1110.905\\,\\an\\ and 1117.161\\,\\an\\ lines. Except for the resonance lines, 1138.076\\,\\an\\ and 1138.210\\,\\an\\ are the strongest isolated S\\,IV lines. interior nucleosynthetic processes, the S\\,IV lines are good diagnostics of temperature.  \\noindent {\\bf Fe\\,III:~} The only isolated Fe\\,III line surviving among the O stars (to O6) is 1091.082\\,\\an.~ However, it is a weak feature. Even at O9.5 it merges with a nearby P\\,IV line in most of our spectra. The Fe\\,III lines are of little help as physical diagnostics in O star spectra.  \\noindent {\\bf Fe\\,IV:~} The dominances of odd iron ion stages, such as Fe$^{3+}$ are narrowly peaked in temperature. As Fig.\\,\\ref{fenums} suggests, the strengths of Fe\\,IV lines peak at subtype B0.2 in main sequence spectra and decrease into the middle-O star spectra.  Nonetheless, several far-UV lines can be utilized as temperature diagnostics. The most useful ones are the isolated, albeit weak, lines at 1005.697\\,\\an\\ and 1156.526\\,\\an.~ The strongest Fe\\,IV feature is 1170.781\\,\\an.~ All Fe\\,IV lines disappear by O6 or O5.  \\noindent {\\bf Fe\\,V:~} This ion contributes by far the largest number of lines to far-UV O star spectra (Fig.\\,\\ref{fenums}). In part because there are so many Fe\\,V lines, few individual features are strong  - even in middle-O stars where the Fe$^{4+}$ ionization fraction peaks. However, one strong unblended line is 1007.292\\,\\an.~  For this case even nearby lines are comparatively weak. Two moderate strength Fe\\,V lines that may provide useful estimates of the iron abundance are 1011.367\\,\\an\\ and 1011.512\\,\\an,~ but these merge in broad lined spectra. The same is true of the Fe\\,V lines at 1018.059\\,\\an\\ and 1018.198\\,\\an.~ Being nearly isolated and of moderate strength, 1043.991\\,\\an\\ is another potentially useful line. The shortest wavelength isolated Fe\\,V lines are 958.288\\,\\an\\ and 958.379\\,\\an.~ These lines are weak but visible across the O4-O9.5 domain. These lines should be used with caution because of the nearby He\\,I line.  \\noindent {\\bf Fe\\,VI:~} This ion represents the most excited ionization state among all iron group elements in our survey. A few isolated lines are visible in O4 to O6 spectra: 1160.561\\,\\an,~ 1170.275\\,\\an,~ and 1186.575\\,\\an.~ As with Fe$^{3+}$, the dominance of Fe$^{5+}$ extends over a relatively small temperature range.  The strengths of these lines begin to weaken in O4 spectra. This trend continues in spectra of the hottest (O2 and O3) stars.  We hope this atlas will serve its intended purposes to aid in the analysis of O star spectra, including population synthesis of star clusters containing massive stars. The author may be contacted for additional information either directly or via MAST."
    },
    "1208/1208.0857_arXiv.txt": {
        "abstract": "The Keck Array (SPUD) began observing the cosmic microwave background's polarization in the winter of 2011 at the South Pole.  The Keck Array follows the success of the predecessor experiments \\bicep\\ and \\bicep2\\cite{Chiang:2009xsa}, using five on-axis refracting telescopes.  These have a combined imaging array of 2500 antenna-coupled TES bolometers read with a SQUID-based time domain multiplexing system.  We will discuss the detector noise and the optimization of the readout.  The achieved sensitivity of the Keck Array is 11.5 $\\mu \\mathrm{K}_{CMB} \\sqrt{s}$ in the 2012 configuration. ",
        "introduction": "\\label{sec:intro}  Inflation, the theory that the universe experienced exponential expansion in its first fraction of a second, was originally introduced as a way to solve the horizon problem--why the universe was nearly homogenous and geometrically flat\\cite{1981PhRvD..23..347G}.  Since then, it predicted many phenomena since confirmed by observation: gaussianity, scale-invariance, and adiabaticity.\\cite{Netterfield:2001yq,Komatsu:2008hk}.  It also is predicted to have produced a gravitational wave background.  The tensor perturbations generated would leave a signature in the curl component of the polarization (B-mode) of the Cosmic Microwave Background radiation (CMB)\\cite{PhysRevLett.78.2058,0004-637X-482-1-6}.  The curl-free component of the polarization (E-mode) is dominated by the scalar, mass density perturbations.  The ratio of the tensor to scalar perturbations ($r$) depends on the energy scale of inflation.  The DASI experiment was the first to detect the polarization of the CMB\\cite{Kovac:2002fg}, detecting the E-mode polarization.  The E-mode polarization power spectrum is determined by the same physics as the temperature power spectrum and thus provided a good test of the CMB paradigm.  The B-mode polarization signal is much smaller than the E-mode, and the current upper limits on $r<0.21$ (95\\% CL) actually come from WMAP and SPT\\cite{Keisler:2011aw} using the temperature information.  BICEP provides the best upper limit from the B-mode power spectrum at $r<0.72$ (95\\% CL)\\cite{Chiang:2009xsa}.  The Keck Array (aka SPUD) is a ground based polarimeter currently observing the CMB.  The Keck Array is the latest of the \\bicep/\\bicep2 family of experiments.  All three experiments use similar optical designs, \\bicep2 uses a focal plane with the same detectors as the Keck Array, and the Keck Array transitioned to a closed-cycle He cooler and increased the number of receivers.  These experiments were designed specifically with the goal of measuring the imprint of gravitational waves from inflation on the polarization of the CMB.  Measuring such a signature will require a lot of sensitivity in large angular scale ($\\ell \\sim 100$) polarization.  Our strategy is to integrate deeply on a 800 $\\mathrm{deg}^2$ patch of sky with a $0.5\\ \\mathrm{deg}$ FWHM beam.  This paper is focused on the Keck Array and its sensitivity.  Several companion papers presented at this conference focus on the current status of \\bicep2 and the Keck Array (Ogburn et al.\\cite{2012SPIE.Ogburn}), the Keck Array optical performance (Vieregg et al.\\cite{2012SPIE.Vieregg}), the performance of the dual planar antennas (O'Brient et al.\\cite{2012SPIE.Obrient}), and the thermal stability performance of \\bicep2 (Kaufman et al.\\cite{2012SPIE.Kaufman}). ",
        "conclusions": "\\end{center} \\end{table}   The optical efficiencies of the detectors are measured in the lab as part of the standard testing regimen.  The measurements are taken by placing a beam-filling, microwave-absorbing cone over the vacuum window of the receiver at room temperature ($\\sim 295$K) and at liquid nitrogen temperature ($\\sim 74$K at the South Pole).  The optical efficiency $\\eta$ is defined as the fraction of power that the bolometer measures compared to total light power input. \\begin{equation} P = \\eta \\int_{\\nu_1}^{\\nu_2} \\frac{h \\nu}{\\exp(h\\nu/kT)-1}\\,\\mathrm{d}\\nu \\label{eq:opeff} \\end{equation} where $P$ is the measured power with the bolometer, and $\\nu_1$ and $\\nu_2$ are the band defining frequencies.  At 150 GHz, both of these temperatures are well in the Rayleigh-Jeans limit ($h\\nu\\ll kT$).  Equation \\ref{eq:opeff} reduces to $\\Delta P = \\eta k_B \\nu \\frac{\\Delta \\nu}{\\nu} \\Delta T$, where $\\frac{\\Delta \\nu}{\\nu}$ is the fractional spectral bandwidth (0.22), $\\nu$ is the band center frequency (145 GHz), and $\\Delta T$ is the difference in temperature of the loads.  From these measurements, the median optical efficiency $\\eta$ is $\\sim 30\\%$.  \\subsection{Readout} \\label{sec:readout}  We apply a voltage bias to the TESs and read out the current from the detectors using time-division multiplexed superconducting quantum interference devices (SQUID)\\cite{deKorte}.  These were developed at NIST and are used in numerous other CMB/submillimeter experiments.  This multiplexing system consists of 3 stages of SQUIDs.  Each detector is inductively coupled to a first stage SQUID (SQ1).  A set of 33 SQ1s are inductively-coupled into a summing coil leading to a single second-stage SQUID (SQ2).  They are multiplexed by simply turning the SQ1 biases on and off.  When its bias is turned off, the SQ1 is superconducting and does not contribute to the summing coil.  The output of the SQ2 then leads to a high gain SQUID series array (SSA) before exiting the cryostat.  The SQ1s and SQ2s are placed on the focal plane unit along with the detectors at 270 mK.  The SSAs are cooled to 4K.  This readout system is also described by Ogburn et al.\\cite{2010SPIE.Ogburn,2012SPIE.Ogburn}.  The electronics for the SQUID multiplexers (Multichannel Electronics or MCE) were developed at UBC\\cite{Battistelli2008}.  These electronics control the current bias and flux feedback (to center the SQUID $V(\\phi)$ modulation curve) for every SQUID.  They also bias the TESs, with only 1 common bias line used per 32 TESs.  The detectors sharing a single bias line are all drawn from the same detector tile and have similar enough properties that this is acceptable.   \\end{center} \\end{table}"
    },
    "1208/1208.5533_arXiv.txt": {
        "abstract": "We studied allowed region on the intergenerational mixing parameters of sleptons from a viewpoint of big-bang nucleosynthesis in a slepton-neutralino coannihilation scenario. In  this scenario, $^7$Li and $^6$Li problems can be solved by considering exotic reactions caused by bound-state effects with a long-lived slepton.  Light element abundances  are calculated as functions of the relic density and lifetime of the slepton which considerably depend on the intergenerational mixing parameters. Compared with observational light element abundances, we obtain allowed regions on the intergenerational mixing. Ratio of selectron component to  stau component, $c_e$, is allowed in $2\\times 10^{-11} \\lesssim c_e \\lesssim 2\\times 10^{-9}$  with solving both the $^7$Li and $^6$Li problems. Similarly, the ratio for smuon, $c_{\\mu}$,  is allowed in $ c_{\\mu} \\lesssim 5\\times 10^{-5}$ for mass  difference between slepton and neutralino, which is smaller than muon mass, and  $ c_{\\mu} \\lesssim 2\\times 10^{-10}$ for the mass difference in range between  muon mass and 125~MeV. We also discuss collider signatures of the slepton decays. We find characteristic double peaks in momentum distribution of event number of the slepton decays with allowed mixing parameters. Discoveries of the double peaks at future collider experiments should confirm our scenario. ",
        "introduction": "%   Long-lived Charged Massive Particles (CHAMP) are predicted in new physics beyond the standard model~\\cite{Fairbairn:2006gg}. Such Long-lived CHAMPs  induce nonstandard nuclear reactions in big-bang nucleosynthesis (BBN), and drastically change  light element abundances~\\cite{Pospelov:2006sc, Kohri:2006cn, Kaplinghat:2006qr, Cyburt:2006uv, Steffen:2006wx,Bird:2007ge, Kawasaki:2007xb, Hamaguchi:2007mp, Jittoh:2007fr, Jedamzik:2007cp, Pradler:2007is, Kawasaki:2008qe, Jittoh:2008eq, Pospelov:2008ta, Kamimura:2008fx, Kusakabe:2008kf, Bailly:2008yy, Bailly:2009pe, Kusakabe:2010cb, Jittoh:2011ni}. By comparing the theoretical predictions with observational data, we can  obtain  constraints  on model parameters relevant with the long-lived CHAMPs. Because the BBN is sensitive to lifetime from $10^{-2}$~sec to $10^{12}$~sec, it should be  one of the best tools to check an existence of the long-lived CHAMPs. \\footnote{See also Refs.~\\cite{Jedamzik:2004er, Kawasaki:2004yh,  Kawasaki:2004qu, Cumberbatch:2007me, Kohri:2008cf, Cyburt:2009pg, Pospelov:2010cw, Kawasaki:2010yh} for long-lived neutral particles.}    The Minimal Supersymmetric Standard Model (MSSM) with R-parity conservation is one of the leading candidate providing long-lived CHAMPs. Several MSSM scenarios predict a long-lived stau as the next lightest  supersymmetric particle (NLSP).  An example is the scenario that the lightest  supersymmetric particle (LSP) is gravitino. In this scenario, the interaction of the gravitino LSP with the stau NLSP is suppressed by Planck scale, and hence the  stau has longevity~\\cite{Ellis:2003dn,Cerdeno:2005eu, Steffen:2006hw,Cyburt:2006uv,Kawasaki:2007xb,Kawasaki:2008qe}. Another example is an axino LSP scenario, in which the decay rate of the stau NLSP is suppressed because of loop processes or a decay constant scale $F_{a}$~\\cite{Covi:2004rb, Freitas:2011fx,Chun:2008rp}. Among such scenarios, the most attractive one is a Bino-like neutralino LSP scenario.  The longevity of the stau NLSP is brought by tight mass degeneracy of the stau and the neutralino LSP.      A remarkable feature of this mass-degenerate scenario is to be free from the $^7$Li  problem \\cite{Cyburt:2008kw}. The $^7$Li  problem is a discrepancy between the observed abundance  of $^7$Li, e.g., ${\\rm Log}_{10}(^7\\text{Li}/\\text{H}) = -9.63 \\pm 0.06$~\\cite{Melendez:2004ni},  and a theoretical one predicted in the standard BBN,  ${\\rm Log}_{10}(^7\\text{Li}/\\text{H}) = -9.35 \\pm 0.06$~\\cite{Jittoh:2011ni}. (There is also a severer observational limit which worsens the fitting~\\cite{Bonifacio:2006au}.)\\footnote{This discrepancy can not be solved even for corrections of corresponding cross sections of nuclear reaction~\\cite{Cyburt:2003ae,Angulo:2005mi}.  Even if we consider nonstandard astrophysical models including diffusion processes~\\cite{Richard:2004pj,Korn:2006tv}, it would be difficult to fit  all of data consistently~\\cite{Lind:2009ta}. See also ~\\cite{Iocco:2012rm} for a recent work related to light element nucleosynthesis in accretion flows. } In order to explain  the observed abundance of the dark matter by the Bino-like neutralino LSP  based on a thermal relic scenario, the so-called stau coannihilation mechanism is required to work well~\\cite{Griest:1990kh}. This mechanism requires of order or smaller than $1\\%$ degeneracy in mass between the neutralino LSP and the stau NLSP. Due to this  tight degeneracy, the stau can be long-lived \\cite{Profumo:2004qt,  Jittoh:2005pq} and trigger exotic nuclear reactions at the BBN era.\\footnote{See also \\cite{Sigurdson:2003vy,Hisano:2006cj, Chuzhoy:2008zy,Borzumati:2008zz,Kohri:2009mi} for another cosmological constraints on CHAMPs.} One of such exotic reactions is an internal conversion process in a bound state of the stau and $^7$Li ($^7$Be) nuclei~\\cite{Jittoh:2007fr, Bird:2007ge}. In the internal conversion process, $^7$Li ($^7$Be) can be destroyed into lighter nuclei and their abundances are fitted to the observational ones. Therefore, in the mass-degenerate scenario, the observed abundances of the dark matter and all light elements are simultaneously realized.      It is further worth examining the BBN in the mass-degenerate scenario. In general, the MSSM has sources of intergenerational mixing, though we had assumed the lepton flavor conservation in our previous works. With the intergenerational mixing, the mass eigenstates are linear combinations of  flavor eigenstates.  As we will discuss later in detail, the intergenerational mixing reduces the number density of the slepton via two types of processes. % Success and failure of solving the $^7$Li problem strongly depends on the number density of the slepton. Furthermore, major and minor of modifications on the light element abundances are also determined by the number density of the slepton. Thus, the modified abundances of the light elements can be predicted by the number density that is a function of mixing parameters. As a result, we can constrain the mixing parameters from the consistency between the modified and the observed abundances of the light elements.         The aim of this work is to predict the slepton mixing parameters in the mass-degenerate scenario. In this scenario, the exotic nuclear reactions catalyzed by the long-lived slepton induce both a  fusion process~\\cite{Pospelov:2006sc} and destruction processes through internal  conversions for $^{7}$Li and $^{7}$Be~\\cite{Jittoh:2007fr,Bird:2007ge,Jittoh:2008eq,Jittoh:2010wh}, and spallations   for $^{4}$He~\\cite{Jittoh:2011ni}. These processes can modify the abundance of the  light elements. In fact, $^6$Li is produced via the former process  after forming a bound state of the slepton with $^4$He.  A fit by the observational abundance $^6$Li/$^7$Li $= 0.046 \\pm 0.022$~\\cite{Asplund:2005yt}   gives an allowed region in a parameter space of the mixing parameters. In addition,  both deuterium (D) and tritium (T) (or $^{3}$He after its decay) are produced nonthermally  by the latter processes.  By adopting recent observational  constraints   D/H = $(2.80 \\pm 0.20) \\times 10^{-5}$~\\cite{Pettini:2008mq} and $^{3}$He/D $<$ 0.87 + 0.27~\\cite{GG03},  an upper bound on the number density of the slepton is obtained as a function of its lifetime. Thus, the upper bound on the number density  results in a lower bounds on the mixing parameters of the slepton. On the other hand, a sufficient number density of the slepton at the bound state formation with $^7$Li ($^7$Be) is required for solving the $^7$Li problem. Then, an upper bound on the mixing parameters is obtained. Hence, intriguingly, the mixing parameters of the slepton are predicted with pinpoint accuracy from both above and below in light of the current observed light element abundances.    This long-lived slepton scenario can be also examined at terrestrial experiments.  For example, due to the longevity longer than {\\it``The first  three minutes of the universe\"},  a small number of the long-lived slepton would  decay in detectors~\\footnote{Collider signatures of the  long-lived slepton in the neutralino LSP scenario is studied in  case that its lifetime is shorter than the beginning of BBN~\\cite{Kaneko:2008re, Kaneko:2011qi} in collider experiments. Measuring the lifetime and the mass difference between the slepton and the neutralino LSP, the intergenerational mixing  will be determined. Combining with the other works on collider signatures of long-lived,  the models can be discriminated.}. Some of them decay into the neutralino LSP and electron (or muon) via the intergenerational mixing. Combined with the predictions on the mixing parameters using the light element abundances, we find monochromatic and diffuse spectrum in momentum distributions of the final electron (or muon). These spectrum are characteristic signatures of the mass-degenerate scenarios, and therefore we can distinguish our scenario from others.       This paper is organized as follows. In the next section, we discuss the relation between the number density of the long-lived slepton and the intergenerational mixing in the mass-degenerate scenarios. Then, we present a set of Boltzmann equations for calculating the number density as a function of the mixing parameters. In Sec.~\\ref{sec:ana}, we analytically estimate bounds on the mixing parameters. Then, in Sec.~\\ref{sec:result}, we show our numerical results. Allowed region on the intergenerational mixing parameter is shown from the viewpoint of observed light element abundances. In Sec.~\\ref{sec:col}, applying the results, we discuss collider implications of this long-lived slepton. Finally, we summarize our results and discuss the prospects. ",
        "conclusions": "\\label{sec:sum} %   We have considered a scenario of the MSSM where the slepton NLSP is long-lived  due to a small mass difference between the NLSP and  the neutralino LSP, and studied the effects of the  intergenerational mixing of sleptons on big-bang  nucleosynthesis.  In this scenario, the so-called internal conversion processes occurs in the bound states between the slepton NLSP and light nuclei in BBN. Then, the $^7$Li and $^6$Li problems can be solved simultaneously when the slepton NLSP is enough  long-lived and its number density is sufficient at the era of BBN. We have analyzed  the yield value and the lifetime of the slepton with the intergenerational mixing as  well as the relic abundances of the light elements including the internal conversion  processes in BBN.   In Sec. III, we have calculated the yield value and the lifetime of the slepton NLSP  and derived upper bounds on the mixing parameters $c_\\mu$.  The upper bounds are obtained by requiring the yield value and the lifetime to  be $Y_{\\tilde{l}} \\sim 2 \\times 10^{-13}$ and $\\tau_{\\tilde{l}} \\ge 10^{3}$ sec. to solve the $^7$Li problem.  In the case of $\\delta m < m_\\mu$, the upper bounds are given as $c_\\mu  \\le 5 \\times 10^{-5}, ~2 \\times 10^{-6}$ and $2 \\times 10^{-7}$ for $\\delta m = 20,~60$ and $100$ MeV, respectively.  On the other hand, in the case of $\\delta m > m_\\mu$, the upper bounds,  $c_\\mu \\le 2 \\times 10^{-10}, ~6 \\times 10^{-11}$ and $2 \\times 10^{-11}$ ,  are given for $\\delta m = 106,~114$ and $122$ MeV, respectively.    We have also analyzed the the relic abundances of the light elements by solving  the Boltzmann equations taking into account all the exotic processes listed in Sec.\\ref{sec:result}. We derived an allowed region on $c_e$ by adopting $2 \\sigma$ uncertainties  for D/H,~$^3$He/D and $^6$Li/$^7$Li, and $2 \\sigma$ and $3 \\sigma$ uncertainty  for $^7$Li/H. We found that, assuming $c_\\mu=0$, the $^7$Li and $^6$Li problems can be solved simultaneously in the range of $2 \\times 10^{-11} \\le  c_e \\le 2 \\times 10^{-9}$. These results are consistent with the estimations  given in Sec. II.    In the end, we have discussed the decays of the slepton in a case of $\\delta m  = 0.07$ GeV and $c_e=2 \\times 10^{-11}$ as an illustrating example.  We found that the $2$-body decay shows a sharp peak at the momentum of  an outgoing electron close to while the $4$-body decay shows a broad peak  in the momentum distribution. The sharp peak is found at the momentum equal  to $\\delta m$, and the expected number of events for $2$-body decay is $100$  times larger than that for $4$-body decay.  These features as well as heavy charged tracks are unique in the  degenerate-mass scenario. In the case that lifetime of the long-lived slepton is  $10^{3}$~sec, even present working detectors can stop the produced long-lived sleptons and accumulate  $\\sim$0.5\\%   of them approximately~\\cite{Asai:2009ka}. If exotic charged tracks of long-lived particle are discovered, new facility may be constructed with higher  accumulation efficiency and high momentum resolution.  Discovery of unique signatures of the decay in addition to charged tracks  of long-lived particle will be a strong evidence of this scenario."
    },
    "1208/1208.2035_arXiv.txt": {
        "abstract": "We present the detection of pulsed gamma-ray emission from the Crab pulsar above 100\\,GeV with the VERITAS array of atmospheric Cherenkov telescopes \\cite{Science}. Gamma-ray emission at theses energies was not expected in pulsar models. The detection of pulsed emission above 100\\,GeV and the absence of an exponential cutoff makes it unlikely that curvature radiation is the primary production mechanism of gamma rays at these energies. ",
        "introduction": "One of the most powerful pulsars in gamma rays is the Crab pulsar \\cite{12,13}, PSR J0534+220, which is the remnant of a historical supernova that was observed in 1054 A.D. It is located at a distance of 6500 light years, has a rotation period of $\\approx$33 ms, a spin-down power of $4.6\\times10^{38}\\,$erg s$^{-1}$ and a surface magnetic field of $3.78\\times10^{12}$\\,G \\cite{14}.  Within the corotating magnetosphere, charged particles are accelerated to relativistic energies and emit non-thermal radiation from radio waves through gamma rays. In general, gamma-ray pulsars exhibit a break in the spectrum between a few hundred MeV and a few GeV. Mapping the cut-off can help to constrain the geometry of the acceleration region, the gamma-ray radiation mechanisms and the attenuation of gamma-rays.  Although past measurements of the Crab pulsar spectrum are consistent with a power law with exponential cut-off, flux measurements above 10\\,GeV are systematically above the best-fit model, suggesting that the spectrum is indeed harder than a power law with exponential cut-off \\cite{13, 16}. However, the statistical uncertainty of the previous data was insufficient to allow a definite conclusion about the spectral shape. In this paper we summarise the recent detection of the Crab pulsar above 100\\,GeV with VERITAS.  \\begin{figure*}[th] \\centering \\includegraphics[width=6.3in]{fig01.eps} \\caption{VERITAS pulse profile of the Crab pulsar at $>120$\\,GeV. The shaded histogram show the VERITAS data. The pulse profile is shown twice for clarity. The dashed horizontal line shows the background level estimated from data in the phase region between 0.43 and 0.94. The data above 100 MeV from the Fermi-LAT \\cite{13} are shown beneath the VERITAS profile. The vertical dashed lines in the panels mark the best-fit peak positions of P1 and P2 in the VERITAS data.  } \\label{profile} \\end{figure*} ",
        "conclusions": "The detection of pulsed gamma-ray emission above 100 GeV provides strong constraints on the gamma-ray radiation mechanisms and the location of the acceleration regions. For example, the shape of the spectrum above the break cannot be attributed to curvature radiation because that would require an exponentially shaped cut-off. In addition, assuming a balance between acceleration gains and radiative losses by curvature radiation, the break in the gamma-ray spectrum is expected to be at $E_{br} = 24\\,$GeV $\\eta^{3/4} \\sqrt{\\xi}$, where $\\eta$ is the acceleration efficiency ($\\eta < 1$) and $\\xi$ is the radius of curvature in units of the light-cylinder radius \\cite{25}. Though $\\xi$ can be larger than one, only with an extremely large radius of curvature would it be possible to produce gamma-ray emission above 100\\,GeV with curvature radiation. It is, therefore, unlikely that curvature radiation is the dominant production mechanism of the observed gamma-ray emission above 100\\,GeV. Two possible interpretations are that either the entire gamma-ray production is dominated by one emission mechanism different from curvature radiation or that a second mechanism becomes dominant above the spectral break energy.  \\bigskip % \\ack This research is supported by grants from the U.S. Department of Energy Office of Science, the U.S. National Science Foundation and the Smithsonian Institution, by NSERC in Canada, by Science Foundation Ireland (SFI 10/RFP/AST2748) and by STFC in the U.K. We acknowledge the excellent work of the technical support staff at the Fred Lawrence Whipple Observatory and at the collaborating institutions in the construction and operation of the instrument. A.\\ N.\\ Otte was in part supported by a Feodor Lynen fellowship from the Alexander von Humboldt Foundation."
    },
    "1208/1208.0801_arXiv.txt": {
        "abstract": "We construct models of universe with a generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n})c^2$ having a linear component and a polytropic component. The linear equation of state $p=\\alpha\\rho c^2$ describes radiation ($\\alpha=1/3$), pressureless matter ($\\alpha=0$), stiff matter ($\\alpha=1$), and vacuum energy ($\\alpha=-1$). The polytropic equation of state $p=k\\rho^{1+1/n} c^2$ may be due to Bose-Einstein condensates with repulsive ($k>0$) or attractive ($k<0$) self-interaction, or have another origin. In this paper, we consider negative indices $n<0$. In that case, the polytropic component dominates in the late universe where the density is low. For $\\alpha=0$, $n=-1$ and $k=-\\rho_{\\Lambda}$, we obtain a model of late universe describing the transition from the matter era to the dark energy era. Coincidentally, we live close to the transition between these two periods, corresponding to $a_2=8.95\\, 10^{25}\\, {\\rm m}$ and $t_2=2.97\\, 10^{17}\\, {\\rm s}$. The universe exists eternally in the future and undergoes an inflationary expansion with the cosmological density $\\rho_{\\Lambda}=7.02\\, 10^{-24}\\, {\\rm g}/{\\rm m}^3$ on a timescale $t_{\\Lambda}=1.46\\, 10^{18} {\\rm s}$. For $\\alpha=0$, $n=-1$ and $k=\\rho_{\\Lambda}$, we obtain a model of cyclic universe appearing and disappearing periodically. If we were living in this universe, it would disappear in about $2.38$ billion years. We make the connection between the early and the late universe and propose a simple equation describing the whole evolution of the universe. This leads to a model of universe that is eternal in past and future without singularity (aioniotic universe). It generalizes the $\\Lambda$CDM model by removing the primordial singularity (Big Bang). This model exhibits a nice ``symmetry'' between an early and late phase of inflation, the cosmological constant in the late universe playing the same role as the Planck constant in the early universe. We interpret the cosmological constant as a fundamental constant of nature describing the ``cosmophysics'' just like the Planck constant describes the microphysics. The Planck density and the cosmological density represent fundamental upper and lower bounds differing by ${122}$ orders of magnitude. The cosmological constant ``problem'' may be a false problem.  We determine the potential of the scalar field (quintessence, tachyon field) corresponding to the generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n})c^2$. We also propose a unification of pre-radiation, radiation and dark energy through the quadratic equation of state $p/c^2=-4\\rho^2/3\\rho_P+\\rho/3-4\\rho_{\\Lambda}/3$. ",
        "introduction": "According to contemporary cosmology, the present energy content of the universe is composed of approximately $5\\%$ ordinary matter, $20\\%$ dark matter and $75\\%$ dark energy \\cite{bt}.  The expansion of the universe began in a tremendous inflationary burst driven by the vacuum energy with the Planck density $\\rho_P=5.16 \\, 10^{99}\\, {\\rm g}/{\\rm m}^3$. Between $10^{-35}$ and $10^{-33}$ seconds after the beginning (Big Bang), the universe expanded by a factor $10^{30}$ \\cite{guth,linde}. Inflation does not offer any explanation for the time before that ``beginning''. The universe then entered in the radiation era and, when the temperature cooled down below approximately $10^3\\, {\\rm K}$, in the matter era \\cite{weinberg}. At present, it undergoes an accelerated expansion \\cite{novae} presumably due to the cosmological constant or to some form of dark energy with negative pressure violating the strong energy condition \\cite{cst}. This corresponds to a second period of inflation, which is different from the first since it is driven by the cosmological energy density $\\rho_{\\Lambda}=7.02\\, 10^{-24}\\, {\\rm g}/{\\rm m}^3$ instead of the Planck density.  Despite the success of this model, the nature of dark matter, dark energy, and of the very early universe (pre-radiation era) remains very mysterious and leads to many speculations.  The phase of inflation in the early universe is usually described by some hypothetical scalar field $\\phi$ with its origin in quantum fluctuations of vacuum \\cite{linde}. This leads to an equation of state $p=-\\rho c^2$, implying a constant energy density,  called the vacuum energy. This energy density is usually identified with the Planck density $\\rho_P$. As a result of the vacuum energy, the universe expands exponentially rapidly on a timescale of the order of the Planck time $t_P=5.39\\, 10^{-44}{\\rm s}$ (early inflation). This phase of inflation is followed by the radiation era described by an equation of state $p=\\rho c^2/3$. In our previous paper (Paper I), we showed that  a unified description of the pre-radiation ($\\rho=\\rho_P$) and radiation ($\\rho\\propto a^{-4}$) eras, connected by  a phase of exponential inflation, could be obtained from a single equation of state of the form $p=(1/3)\\rho(1-4\\rho/\\rho_P)c^2$. This provides a non-singular model of the early universe.   The phase of acceleration in the present universe is usually ascribed to the cosmological constant $\\Lambda$  which is equivalent to a constant energy density $\\rho_{\\Lambda}=\\Lambda/(8\\pi G)$ called dark energy. This can be modeled by an equation of state $p=-\\rho c^2$, implying a constant energy density identified with the cosmological density  $\\rho_{\\Lambda}$. As a result of the dark energy, the universe expands exponentially rapidly on a  timescale of the order of the cosmological time $t_{\\Lambda}=1.46\\, 10^{18} {\\rm s}$ (de Sitter solution). This leads to a second phase of inflation (late inflation). Inspired by the analogy with the early inflation, some authors have represented the dark energy by a scalar field called quintessence \\cite{quintessence}. As an alternative to the quintessence, other authors  \\cite{chaplygin} have proposed to model the acceleration of the universe by an exotic fluid with an equation of state of the form $p=-A/\\rho$  called the Chaplygin gas (see \\cite{cst} for a complete list of references). At late times, this equation of state leads to a constant energy density implying  an exponential growth of the scale factor that is similar to the effect of the cosmological constant. At earlier times, this equation of state returns the results of the cold dark matter model. Therefore, it provides a unification of dark matter ($\\rho\\propto a^{-3}$) and dark energy ($\\rho=\\rho_{\\Lambda}$) in the late universe. Furthermore, it gives a real velocity of sound which is non-trivial for fluids with negative pressure.  Some generalizations of this equation of state have been considered in the form $p=-A/\\rho^{a}$ with $a\\ge -1$ \\cite{chaplygin,cst}. As mentioned in \\cite{chaplygin}, the Chaplygin gas has some connection with string theory and can be obtained from the Nambu-Goto action for $d$-branes moving in a $(d+2$)-dimensional spacetime in the light-cone parametrization. Furthermore, it is the only fluid which, up to now, admits a supersymmetric generalization.  From a theoretical point of view, it is desirable to study models of universe with a generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n}) c^2$ having a standard linear component and a polytropic component. The linear equation of state $p=\\alpha\\rho c^2$ describes radiation ($\\alpha=1/3$), pressureless matter ($\\alpha=0$), stiff matter ($\\alpha=1$), and vacuum energy ($\\alpha=-1$). The polytropic equation of state $p=k\\rho^{\\gamma} c^2$ with $\\gamma=1+1/n$ may be due to Bose-Einstein condensates with repulsive ($k>0$) or  attractive ($k<0$) self-interaction \\cite{c4}, or have another origin. When $n>0$, the polytropic component dominates the linear component when the density is high. These models, studied in Paper I, describe the early universe. Conversely, when $n<0$, the polytropic component dominates  the linear component when the density is low. These models, studied in the present paper, describe the late universe. As we shall see, these two studies are strikingly symmetric.  Interestingly, this symmetry seems to reflect the true evolution of the universe.    In this paper, we propose an exhaustive study of the equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n}) c^2$ with $n<0$. When $k>0$, the universe exhibits a future peculiarity: Its density vanishes in infinite time ($n\\le -2$), or periodically in time ($n>-2$), while its radius tends to a constant value (the universe is singular when $n>-1$ since the pressure becomes infinite in finite time).  For $\\alpha=0$, $n=-1$ and $k=\\rho_{\\Lambda}$, we obtain a model of cyclic universe, equivalent  to the anti-$\\Lambda$CDM model, in which the universe ``disappears'' (it becomes empty) and ``reappears'' periodically. According to this model, the universe would disappear in about $2.38$ billion years.   When $k<0$, the universe exists for all times in the future and there is no singularity. For $\\alpha=0$, $n=-1$ and $k=-\\rho_{\\Lambda}$, we obtain a model of late universe, equivalent to the $\\Lambda$CDM model, describing in a unified manner the transition from the matter era to the dark energy era. Coincidentally, we live close to the transition between these two periods, corresponding to $a_2=8.95\\, 10^{25}\\, {\\rm m}$ and $t_2=2.97\\, 10^{17}\\, {\\rm s}$. This universe exists eternally in the future and undergoes an inflationary expansion with the cosmological density $\\rho_{\\Lambda}=7.02\\, 10^{-24}\\, {\\rm g}/{\\rm m}^3$ on a timescale $t_{\\Lambda}=1.46\\, 10^{18} {\\rm s}$.     The paper is organized as follows. In Sec. \\ref{sec_basic}, we recall the basic equations of cosmology. In Secs. \\ref{sec_ges} and \\ref{sec_dark}, we study the generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n}) c^2$ for any value of the parameters $\\alpha$, $k$ and $n<0$. In Sec. \\ref{sec_solid}, we consider the case $\\alpha=0$, $n=-1$, and $k=- \\rho_{\\Lambda}$, providing a model of non-singular inflationary universe describing the transition from the matter era to the dark energy era ($\\Lambda$CDM model). In Sec. \\ref{sec_cyclic}, we consider the case $\\alpha=0$, $n=-1$, and $k= \\rho_{\\Lambda}$, leading to a model of peculiar cyclic universe (anti-$\\Lambda$CDM model).  In Sec. \\ref{sec_simple}, we discuss the connection between the early and the late universe, and propose a simple equation [see Eq. (\\ref{pr7})] describing the whole evolution of the universe. This leads to a model of universe that is eternal in past and future without singularity (aioniotic universe). This model exhibits a nice ``symmetry'' between an early and late phase of inflation, the cosmological constant in the late universe playing the same role as the Planck constant in the early universe. This model refines the standard $\\Lambda$CDM model by removing the primordial singularity. In Sec. \\ref{sec_scalar}, we determine the potential of the effective scalar field (quintessence, tachyon field) corresponding to the generalized equation of state $p=(\\alpha \\rho+k\\rho^{1+1/n}) c^2$. ",
        "conclusions": "In this paper, and in the previous one, we have carried out an exhaustive study of the generalized equation of state (\\ref{b10}), considering all the possible cases. We have obtained the following results (we assume here $\\alpha\\neq -1$):  (i) For $n>0$ and $k<0$, the universe undergoes an early inflation. It starts from $t=-\\infty$ with a vanishing radius and a finite density $\\rho_{max}$. Its radius increases with time while its density decreases. The universe exists at any time in the past and is non-singular.  (ii) For $n>0$ and $k>0$, the universe undergoes a new type of primordial singularity. It starts at $t=0$ with a finite radius and an infinite density. For $t>0$, its radius increases while its density decreases.  (iii) For $n<0$ and $k<0$, the universe undergoes a late inflation. Its radius increases to $+\\infty$ as  $t\\rightarrow +\\infty$ while its density decreases to a constant value $\\rho_{min}$. The universe exists at any time in the future and is non-singular.  (iv) For $n<0$ and $k>0$, the universe undergoes a future peculiarity. Its radius increases to a maximum value while its density decreases to zero (the universe ``disappears''). For $n\\le -2$, this peculiarity is reached in infinite time. For $n>-2$, this peculiarity is reached in a finite time (and there is a future singularity when $n>-1$ due to the divergence of the pressure). Then, the radius decreases to zero while the density increases to $+\\infty$. These phases of expansion and contraction continue periodically (cyclic universe).   The generalized equation of state (\\ref{b10}) can be viewed as a ``mixture'' of a linear equation of state describing a classical universe filled with radiation ($\\alpha=1/3$) or matter ($\\alpha=0$), and a polytropic equation of state whose origin remains to be understood (in Paper I, we have mentioned the connection with Bose-Einstein condensates, but other possibilities may be contemplated).  Positive indices describe the early universe (the polytropic component dominates the linear component when the density is high) and negative indices describe the late universe (the polytropic component dominates the linear component when the density is low). A positive polytropic pressure ($k>0$) leads to singular or peculiar models, while a negative polytropic pressure ($k<0$) leads to non-singular models that exist for all times.   For $k<0$, the generalized equation of state (\\ref{b10}) implies the existence of upper and lower bounds on the density. It is natural to identify the upper bound with the Planck density $\\rho_{max}=\\rho_P$ (quantum mechanics) and the lower bound with the cosmological density $\\rho_{min}=\\rho_{\\Lambda}$ (general relativity). By taking $n=1$ in the early universe and $n=-1$ in the late universe, we have obtained a non-singular model that is consistent with the known properties of our universe. This model improves the standard $\\Lambda$CDM model by removing the primordial singularity. An attractive feature of this model is its simplicity and the  ``symmetry'' that it reveals between the past and the future. Furthermore, this model admits a scalar field representation based on a quintessence field (in the early and late universe) or a  tachyon field (in  the late universe).   By using only the principle of ``simplicity'', and rejecting singular or peculiar solutions, we have obtained in a natural manner a simple model of universe. It corresponds to an ``aioniotic'' universe that exists eternally in past and future and whose early and late evolutions present some striking symmetry. The remarkable point is that this model, which is obtained in a purely theoretical manner independent of observations, turns out to be fully consistent with the known structure of the real universe. A theoretical challenge would now be to {\\it justify} a quadratic equation of state $p=-4\\rho^2c^2/(3\\rho_P)$ ($n=1$) in the early universe and a constant pressure $p=-\\rho_{\\Lambda}c^2$ ($n=-1$) in the late universe. It may also be interesting to study an equation of state of the form $p/c^2=-(\\alpha+1)\\rho^2/\\rho_P+\\alpha\\rho-(\\alpha+1)\\rho_{\\Lambda}$ exhibiting two phases of inflation \\cite{prep}. For $\\alpha=1/3$, it provides a unification of pre-radiation, radiation, and dark energy (see Appendix \\ref{sec_unification}).     Our study suggests that the Planck density $\\rho_P$ and the cosmological density $\\rho_{\\Lambda}$  represent fundamental upper and lower bounds. These bounds are responsible for a phase of early and late inflation. It is oftentimes argued that the dark matter density $\\rho_{\\Lambda}$ should be identified with the vacuum energy. Since the vacuum energy is of the order of the Planck density $\\rho_P$, which is $10^{122}$ times larger than the cosmological density, this leads to the so-called cosmological constant problem \\cite{weinbergcosmo}. Actually, if the Planck density and the cosmological density represent fundamental upper and lower bounds, it is not surprising that they differ by about ${122}$ orders of magnitudes. Therefore, the origin of the cosmological density should not be sought in quantum mechanics but in a new theory of cosmophysics based on general relativity.  We have also studied singular models of universe corresponding to a positive polytropic pressure ($k<0$). Although these models do not appear to be selected by nature, they are interesting on a mathematical point of view. For $n>0$, we have obtained a model exhibiting a primordial singularity at $t=0$ at which the radius has a finite  value while the density and the temperature are infinite. For $n<0$, we have obtained models exhibiting future peculiarities in which the density of the universe vanishes while its radius reaches a maximum value. This peculiarity may occur in infinite time ($n\\le -2$), or periodically in time ($n>-2$). In particular, we have obtained a simple analytical solution of a cyclic universe (anti-$\\Lambda$CDM model) in which the density disappears and re-appears periodically. Since this solution is always decelerating (in the phase of expansion), this peculiar model is not compatible with the known properties of our universe.  There remains a type of solutions that we have not described so far: The case where the  density increases as the universe expands ($w<-1$). These {\\it a priori} un-natural solutions which violate the null dominant energy condition  are associated with phantom scalar fields \\cite{caldwell,caldwellprl}. They are studied in Paper III in which we construct models of ``phantom universes''. Actually, there is a rich recent literature \\cite{cst,ghosts} on these solutions since observations do not exclude the possibility that we live in a phantom universe \\cite{observations}. These models usually predict a future singularity at which the radius and the density of the universe become infinite in a finite time. This would lead to the death of the universe in a Big Rip or a Cosmic Doomsday \\cite{caldwellprl}.  The main feature of the generalized equation of state (\\ref{b10}) is its polytropic component $p=k\\rho^{\\gamma}c^2$. A polytropic equation of state occurs in many situations of astrophysical interest \\cite{bt}. For example, a polytropic equation of state with index $\\gamma=5/3$ describes adiabatic processes in main sequence stars \\cite{chandra}. Compact stars such as white dwarfs, neutron stars and BEC stars are also described by a polytropic equation of state. Classical  white dwarf stars correspond to a polytropic  index $n=3/2$ and relativistic white dwarf stars correspond to a polytropic  index $n=3$ \\cite{chandra,st}. Newtonian and semi-relativistic BEC stars with a self-interaction correspond to a polytropic index $n=1$ \\cite{ch}. From a mathematical point of view, polytropic stars represent a particular class of steady states of the Euler-Poisson system.  Similarly, stellar polytropes  represent a particular class of steady states of the  Vlasov-Poisson system \\cite{bt}.  Polytropic distributions also appear in statistical physics, in relation with generalized thermodynamics \\cite{tsallisbook} and nonlinear Fokker-Planck equations \\cite{frank}. In previous works, we have studied polytropic equations of state in several situations of astrophysical \\cite{aaantonov,cs}, physical \\cite{cc}, and biological \\cite{nfp} interest. It was therefore natural to consider the application of polytropic equations of state in cosmology also.     \\appendix"
    },
    "1208/1208.0340_arXiv.txt": {
        "abstract": "Direct $N$-body simulations of star clusters in a realistic Milky Way-like potential are carried out using the code \\texttt{NBODY6}. Based on these simulations a new relationship between scale size and galactocentric distance is derived: the scale size of star clusters is proportional to the hyperbolic tangent of the galactocentric distance. The half-mass radius of star clusters increases systematically with galactocentric distance but levels off when star clusters orbit the galaxy beyond $\\sim$40 kpc. These simulations show that the half-mass radius of individual star clusters varies significantly as they evolve over a Hubble time, more so for clusters with shorter relaxation times, and remains constant through several relaxation times only in certain situations when expansion driven by the internal dynamics of the star cluster and the influence of the host galaxy tidal field balance each other. Indeed, the radius of a star cluster  evolving within the inner 20 kpc of a realistic galactic gravitational potential is severely truncated by tidal interactions and does not remain constant over a Hubble time. Furthermore, the half-mass radius of star clusters measured with present day observations bears no memory of the original cluster size. Stellar evolution and tidal stripping are the two competing physical mechanisms that determine the present day size of globular clusters. These simulations also show that extended star clusters can form at large galactocentric distances while remaining fully bound to the host galaxy. There is thus no need to invoke accretion from an external galaxy to explain the presence of extended clusters at large galactocentric distances in a Milky Way-type galaxy. ",
        "introduction": "Star clusters are an increasingly diverse family. During the last decade the discovery of stellar systems brighter, larger, and more massive than the ``standard\" star cluster has blurred the distinction between globular clusters and dwarf galaxies. Ultra-compact dwarfs (UCDs) are the prime example of stellar systems with physical parameters between those of globular clusters and dwarf elliptical galaxies (Hilker et al.\\ 1999; Drinkwater et al.\\ 2000). At the other end of the luminosity range, faint and extended star clusters have also been discovered, notably as satellites of the Andromeda galaxy (Huxor et al.\\ 2005; Mackey et al.\\ 2006) and NGC 1023 (Larsen \\& Brodie 2000).  These newly discovered stellar systems have bridged the gap in physical size thought to exist between globular clusters and compact elliptical galaxies (Gilmore et al.\\ 2007). The well defined linear relations between physical size and total magnitude for galaxies with masses greater than $10^8 M_{\\odot}$ can be extrapolated to UCDs. However, dwarf galaxies and star clusters form two branches in the size-magnitude plane where these two physical parameters are uncorrelated (Misgeld \\& Hilker 2011 their Figure 1; see also McLaughlin 2000).  In recent years, studies based on the relatively wide field of view and superb resolution of the Advanced Camera for Surveys onboard the {\\it Hubble Space Telescope (HST)}, have found a strikingly constant median effective radius, or equivalently half-light radius, for extragalactic globular clusters of $<r_h>\\sim 3$ pc. Jord\\'an et al.\\ (2005) provide a prime example of such work as they accurately determined the structural parameters of thousands of globular clusters associated with 100 early type galaxies of the Virgo cluster (ACS Virgo Cluster Survey). Masters et al.\\ (2010) replicated this work for 43 galaxies in the Fornax cluster (ACS Fornax Cluster Survey). One of the main findings of these papers is that thousands of star clusters spanning more than four magnitudes in luminosity have the same median value of $<r_h> \\sim 3$ pc.  Why is the size distribution of star clusters narrowly centered around three parsecs? Why do only a few clusters become extended with effective radii of ten parsecs or more? In this work, advanced $N$-body models are carried out with the aim of determining the most important physical mechanisms that mold the characteristic radii of star clusters. The impact of the host galaxy tidal field on the size of orbiting star clusters is probed in detail by evolving several models at different galactocentric distances ($R_{GC}$).  The empirical qualitative dependence between size and galactocentric distance of star clusters has been clearly established in several observational studies beginning with the work of Hodge (1960, 1962). The $N$-body models that have been performed allow us to quantify the influence of the tidal field, generated by a Milky-Way or M31 type galaxy, on satellite star clusters. We thus determine a new relation between the scale sizes of star clusters and galactocentric distance. This relation is a proxy for the host galaxy gravitational potential. ",
        "conclusions": ""
    },
    "1208/1208.5279_arXiv.txt": {
        "abstract": "We study dynamical capture binary neutron star mergers as may arise in dense stellar regions such as globular clusters.  Using general-relativistic hydrodynamics, we find that these mergers can result in the prompt collapse to a black hole or in the formation of a hypermassive neutron star, depending not only on the neutron star equation of state but also on impact parameter.  We also find that these mergers can produce accretion disks of up to a tenth of a solar mass and unbound ejected material of up to a few percent of a solar mass. We comment on the gravitational radiation and electromagnetic transients that these sources may produce. ",
        "introduction": "Merging binary neutron stars (NSs) warrant detailed study because these systems promise to be rich sources of both gravitational and electromagnetic (EM) radiation, probing strong-field gravity and nuclear density physics.  NS--NS mergers are a primary source targeted by gravitational wave (GW) detectors~\\citep[such as LIGO;][]{LIGO}.  They are also candidates for short gamma-ray burst (SGRB) progenitors and several other EM counterparts~\\citep{2012ApJ...746...48M,2041-8205-736-1-L21,2012arXiv1204.6242P} which could potentially be observed by current and upcoming wide-field survey telescopes like PTF~\\citep{2009PASP..121.1334R}, Pan-STARRS~\\citep{2004SPIE.5489...11K}, and LSST~\\citep{2009arXiv0912.0201L}.  There have been numerous studies of primordial binary NS mergers (see e.g.~\\cite{faber_review}), which will have essentially zero orbital eccentricity when they enter the frequency band of ground-based GW detectors. However, binaries may also form via $n$-body interactions in dense stellar regions and some fraction of them will have sizable eccentricity at merger. A likely environment to find such binaries is globular clusters (GCs) that have undergone core collapse \\citep{fabian75,grindlay2006}. In~\\cite{lee2010}, it was argued that the rate of tidal capture and collision of two NSs in GCs (using M15 as a prototype) peaked around $z\\simeq0.7$ at values of $\\sim50$ yr$^{-1}$ Gpc$^{-3}$ (falling to $\\sim30$ yr$^{-1}$ Gpc$^{-3}$ by $z=0$) and was consistent with the observed SGRB rate. However, this does not take into account natal kicks. A recent simulation of M15 that assumed a modest NS retention fraction of $5\\%$ found $\\sim1/4$ fewer NSs in the central $0.2$ pc~\\citep{2011ApJ...732...67M} compared to an earlier study that ignored kicks~\\citep{dull,2003ApJ...585..598D}, implying the above rate (scaling as the number density squared) could be overestimated by an order of magnitude. On the other hand, observations suggest that the NS retention fraction in some GCs can be as large as $20\\%$~\\citep{2002ApJ...573..283P}, and the model of~\\cite{lee2010} did not take into account other channels that could lead to binary merger within a Hubble time, such as Kozai resonance in a triple system~\\citep{Thompson:2010dp}.  The above discussion focused on GC environments, and similar interactions in galactic nuclei would also add to the rates~\\citep{oleary,Kocsis:2011jy,Antonini2012}. Still, it is far from certain that high eccentricity mergers occur frequently enough to expect observation with the upcoming generation of GW detectors. However, it is also not implausible that they do, and as eccentric NS mergers may also produce distinguishable EM emission compared to quasi-circular mergers, it behooves us to understand both systems from a multi-messenger perspective.   In~\\cite{ebhns_letter} and \\cite{bhns_astro_paper}, black-hole--neutron-star (BH--NS) mergers formed through dynamical capture were found to exhibit a rich variation with impact parameter, in some cases producing sizable disks and amounts of unbound material. In~\\cite{gold}, several eccentric NS--NS mergers were studied using a $\\Gamma=2$ equation of state (EOS) and shown to exhibit $f$-mode excitation during close encounters. There have also been studies of BH--NS and NS--NS collisions with Newtonian gravity~\\citep{lee2010,2012arXiv1204.6240R} showing similar variation in the outcomes.  In this Letter, we study dynamical capture NS--NS mergers for a range of impact parameters using general-relativistic hydrodynamics (GRHD). We also consider several different NS EOSs because of the uncertainty regarding the correct description of matter above nuclear densities. One of the important issues we address for the first time is if these mergers can produce hypermassive neutron stars (HMNSs).  In studies of quasi-circular systems it was found that thermal energy from the merger, as well as differential rotation, could support long-lived HMNSs for some EOSs~\\citep[e.g.,][]{PhysRevLett.107.051102} and that this would be imprinted in the GW signal and resulting disk properties.  HMNSs with longer lifetimes can also build up significant magnetic fields which can power strong EM transients during the collapse to a BH~\\citep{2011arXiv1112.2622L}. For dynamical capture binaries, the amount of angular momentum, and likely the amount of shock heating, will be strong functions of impact parameter, suggesting HMNS formation will be as well.  Another notable feature of dynamical capture NS--NS mergers is their potential to produce unbound nuclear material which will decompress and form heavy nuclei via the $r$-process ~\\citep{1974ApJ...192L.145L,Rosswog:1998gc,Li:1998bw}; subsequent radioactive decay could produce observable emission. Recent work~\\citep{Fischer,Arcones} suggests processes like NS--NS mergers may be needed to supplement the supernovae $r$-process yield in accounting for the observed abundances. Though simulations of quasi-circular NS--NS mergers using Newtonian or conformally flat gravity have found suitable ejecta, they seem to be in tension with fully general-relativistic results which find negligible amounts of ejecta~\\citep{faber_review}. This is {\\em arguably} because of strong-field GR effects, such as BH formation and the existence of innermost stable orbits. As we show, dynamical capture mergers are more promising sources of ejecta, presumably as the stars are less bound when disruption occurs.  In the remainder of this Letter, we outline our methods for simulating NS--NS mergers with GRHD, discuss the merger dynamics for a range of impact parameters and three different EOSs, and comment on potential GW and EM counterparts.  We find that, while the GW signals from these mergers may be challenging to detect with upcoming ground-based detectors, they have the potential to source numerous EM transients. Non-merging close encounters can induce tidal deformations strong enough to crack the NSs' crusts; a merger where the total mass is above the maximum mass of a single NS can either promptly collapse to a BH or produce a hot, rapidly rotating HMNS, where the latter outcome tends to have more massive disks and ejected material. ",
        "conclusions": "We have performed GRHD simulations modeling dynamical capture NS--NS mergers, giving direct estimates of the corresponding GW emission and merger outcome varying impact parameter, EOS, and mass ratio. By measuring pre-merger tidal deformation and post-merger stripped material (bound and ejected), we have also speculated on related EM transients.  Regarding transients that may precede the merger, non-merging close encounters can lead to tidal deformations strong enough to crack the NSs' crust and tap into the $\\sim10^{46}$ erg stored in elastic energy~\\citep{1995MNRAS.275..255T}, potentially causing flaring activity from milliseconds up to possibly a few seconds before merger. Though a different mechanism and time scales, the signature could be similar to resonance induced cracking for quasi-circular inspirals proposed in~\\cite{PhysRevLett.108.011102}. The cracking of the NS crust is one possible explanation for SGRB precursors observed by \\emph{Swift}~\\citep{precursors}.  We find that dynamical capture mergers can result in prompt BH formation or the formation of an HMNS depending on impact parameter and EOS.  The HMNSs will be long lived due to their rapid rotation and thermal energy, giving them the potential to seed large magnetic fields and source intense transients during collapse.  In contrast to what was found in general-relativistic studies of quasi-circular NS--NS mergers, we find that dynamical capture mergers can result in massive disks even for equal mass binaries, and can result in up to a few percent of a solar mass in ejecta.  This mildly relativistic ejecta can produce potentially observable optical and radio transients. The amount of ejecta found here is similar to the $0.009$--$0.06$ $M_{\\odot}$ found with Newtonian gravity~\\citep{2012arXiv1204.6240R}, though not for comparable impact parameters ($r_p\\leq5$). However, what is qualitatively consistent with the Newtonian setups is that we observe the largest amounts of unbound material for grazing collisions.  Regarding GW detectability, the high frequency of the merger-ringdown or quasi-periodic signals from the HMNS will be difficult to observe with AdLIGO.  Individual bursts from close encounters would also not be detectable except for very nearby events. For example, an $r_p=10$, HB EOS merger at $d=100$ Mpc has sky-averaged S/N for AdLIGO of $\\approx0.9$. This implies that if dynamical capture NS--NS mergers constitute a fraction of SGRB progenitors, a further subset of these will not have a detectable GW counterpart. GW signals from larger $r_p$ binaries undergoing numerous close encounters would have larger S/N, and the timing between bursts will be a sensitive function of the orbital energy, containing information about the EOS, for example. We defer a detailed study of GW detectability to future work~\\citep{upcoming_ecc_gwave}."
    },
    "1208/1208.1941_arXiv.txt": {
        "abstract": "The radio structures and optical identifications of a sample of 242 sources classified as double-double radio sources by Proctor (2011) from a morphological study of sources in the FIRST (Faint Images of the Radio Sky at Twenty centimeters) survey (2003 April release, 811,117 entries) have been examined. We have been able to confirm only 23 of these as likely to be double-double radio galaxies (DDRGs), whose structures could be attributed to episodic nuclear activity in their host galaxies. A further 63 require either higher-resolution radio observations or optical identifications to determine whether these are DDRGs. The remaining sources are unlikely to be DDRGs. We have examined the luminosities, sizes and symmetry parameters of the DDRGs and the constraints these place on our understanding of these sources. ",
        "introduction": "\\label{sec:Intro} An important and interesting question in our understanding of active galactic nuclei (AGN) is whether their nuclear activity is usually episodic in nature. There have been suggestions that black holes grow during AGN phases with the total life time of the active phases ranging from $\\sim1.5\\times10^8$ to $10^9$ yr (e.g. Marconi et al. 2004). Recurrent activity could also have significant implications in feedback processes in active galaxies and on the evolution of galaxies themselves. Although there has been increasing evidence of recurrent activity in AGN from both radio and X-ray observations (see Saikia \\& Jamrozy 2009 for a review), it is not clear how ubiquitous this phenomenon might be. For example from deep low-frequency observations Sirothia et al. (2009) did not find any unambiguous evidence of recurrent activity in a sample of 374 small- sized sources. Marecki (2012) has interpreted the structure of the highly asymmetric giant radio galaxy J1211+743 (Pirya et al. 2011) to be due to recurrent nuclear activity. To understand these varied aspects as well as the range of time scales of episodic activity and possible reasons for it (see Kaiser, Schoenmakers \\& R\\\"ottgering 2000; Czerny et al. 2009; Brocksopp et al. 2011), it is necessary to significantly increase the number of sources with evidence of recurrent activity beyond the couple of dozen or so that are presently known (see Saikia \\& Jamrozy 2009). At present, there appears to be a reasonably wide range of time scales of episodic activity which has been estimated from spectral and dynamical ages of the outer and inner lobes of a few DDRGs. These range from $\\sim10^5$ yr for 3C293 (Joshi et al. 2011) to $\\sim10^8$ yr for the largest Mpc-scale DDRGs (e.g. Schoenmakers et al. 2000; Konar et al. 2006).  Approximately 10 per cent of AGN are luminous at radio wavelengths, referred to as radio- loud objects. Drawing an analogy with X-ray binary systems in our Galaxy, a number of authors have suggested that radio activity in AGN could itself be episodic in nature (Nipoti, Blundell \\& Binney 2005; K\\\"ording, Jester \\& Fender 2006). One of the clearest signs of episodic activity is seen in radio-loud objects where there are more than one pair of distinct outer lobes, which can be unambiguously ascribed to different cycles of activity. Although most of these sources exhibit two cycles of activity, and are referred to as DDRGs, there are two examples of sources which appear to possibly exhibit three cycles of activity, namely B0925$+$420 (Brocksopp et al. 2007) and J140948.85$-$030232.5 (Hota et al. 2011), the latter being associated with a spiral host galaxy. Almost all the sources exhibiting evidence of a double-double structure are associated with galaxies, with the possible exception of 4C02.27 which is associated with a quasar (Jamrozy, Saikia \\& Konar 2009). Evidence of episodic activity has also been reported from a combination of X-ray and radio observations, where inverse-Compton scattered X-ray emission from old electrons in the relic lobes are seen along with synchrotron emission from the recent cycle of activity. Sources where this has been reported include 3C191 and 3C294 (Erlund et al. 2006), as well as the well-studied object Cygnus A (Steenbrugge, Blundell \\& Duffy 2008; Steenbrugge, Heywood \\& Blundell 2010).  To understand the nature of these sources and possible reasons for their episodic activity, we need to enlarge the sample of objects. As a first step we have focused on the FIRST survey (Becker, White \\& Helfand 1995), where Proctor (2011) has done a classification of the structures of sources into different categories. Of interest here are the 242 sources Proctor (2011) has classified as DDRGs based on the identification of at least four different components in the radio images. Since the mere existence of four components does not guarantee the source to be a DDRG, we have examined the radio structure as well as the optical identifications of all the 242 sources to identify those we believe to be good examples of DDRGs, possible examples which require further observations and also sources which do not appear to be DDRGs. We describe briefly the methodology we have adopted in Section 2, describe the DDRGs identified from this survey in Section 3, and discuss the nature of these sources in Section 4. ",
        "conclusions": "\\label{sec:con_remk} Early studies of DDRGs suggested that these are likely to be associated with giant radio sources, yielding time scales of episodic activity of $\\sim$10$^8$ yr or so (e.g. Schoenmakers et al. 2001; Saikia, Konar \\& Kulkarni 2006; Saikia \\& Jamrozy 2009 and references therein). Identification of 23 DDRGs from the FIRST survey has shown that these often occur in sources with overall projected sizes of hundreds of kpc as is seen in the misaligned DDRG 3C293 (Joshi et al. 2011). Salter et al. (2010) also speculated that radio emission from the archetypal GPS source CTA21 may be episodic. Examples of DDRGs appear to occur over a wide range of size scales, and surveys of different resolutions are required to be able to identify these objects. Detailed spectral and dynamical age studies will help us explore the range of time scales of episodic nuclear activity. In this paper we also list the candidate DDRGs which require further observations to determine whether it is a DDRG, and also the ones we classify as non-DDRGs and WATs in the Appendix."
    },
    "1208/1208.5490_arXiv.txt": {
        "abstract": "We study primordial gravitational waves (PGWs)  in the Horava-Lifshitz (HL) theory of quantum gravity, in which high-order spatial derivative operators, including the ones violating parity, generically appear  in order for the theory to be power-counting renormalizable and ultraviolet (UV) complete. Because of both parity violation and non-adiabatic evolution of the modes due to a modified dispersion relationship, a large polarization of PGWs becomes possible, and it could be well within the range of detection of the BB, TB and EB power spectra of  the forthcoming  cosmic microwave background (CMB) observations. ",
        "introduction": "With the arrival of the precision era of cosmological measurements,   temperature and polarization maps of  CMB  with unprecedented accuracy  will soon become available \\cite{KDM}, and shall provide a wealth of data concerning the physics of the early universe, including inflation \\cite{Guth}, a dominant paradigm, according to which primordial  density and PGWs  were created from quantum  fluctuations in the very early universe. The former  grows to produce the observed  large-scale  structure, and meanwhile creates CMB temperature anisotropy, which was already detected by the Cosmic Background Explorer (COBE) almost two decades  ago \\cite{COBE}. PGWs, on the other hand,  produce not only  temperature anisotropy, but also a distinguishable  signature in CMB polarization. In particular, decomposing the polarization into two modes: one is curl-free, the E-mode, and the other is divergence-free, the B-mode, one finds that the B mode pattern cannot be produced by density fluctuations. Thus, its detection would provide a unique signature for the existence of PGWs \\cite{seljak}. PGWs normally produce the TT, EE, BB and TE spectra of CMB, but the spectra of TB and EB vanish when  the parity of the PGWs  is conserved.  However, if the theory is chiral, the power spectra of right-hand and left-hand PGWs can have different amplitudes, and then induce non-vanishing TB and EB correlation in  large scales \\cite{lue}. This provides the opportunity to directly detect the chiral asymmetry of the theory by observations  \\cite{lue,saito,kamionkowski}.  With the above  motivations, the studies of PGWs  have attracted a great deal of attention, and various aspects have been explored \\cite{KDM}. Current 7-year Wilkinson Microwave Anisotropy Probe (WMAP) observations give the constraint on the tensor-to-scalar ratio $r<0.36$\\cite{wmap7}, and the 9-year data give the similar result $r<0.38$ \\cite{wmap9}. However, if combining with other cosmological observations, the recent 9-year WMAP gives the tightest constraint $r<0.13$ at $95\\%$ confidence level \\cite{wmap9}, which corresponds to the amplitude of  the PGWs $\\Delta_h^2<3.03\\times10^{-10}$. It should be noted that  they impose no constraint on their chirality \\cite{saito}.   In this paper, we investigate  the possibility to detect the  chirality  of PGWs through  three information channels, BB, TB and EB of CMB, in the recently proposed  HL theory of quantum gravity, in which parity violation happens generically \\cite{horava}. Such a detection    places the theory  directly under tests, and provides  a smoking gun in its parity violation  in the early universe. This   represents   one of few  observations/experiments that one can construct currently as well as  in the near future, considering the quantum nature of the theory.  The HL theory  is power-counting renormalizable, because of the presence  of high-order spatial derivative operators. The exclusion of high-order time derivative operators, on the other hand,  renders  the theory unitary, whereby it  is expected to be UV complete. In the infrared (IR), the low-order derivative operators take over [cf. Eq.(\\ref{DFE})] and presumably provides a healthy IR limit \\cite{horava}. When applying it to cosmology, various  remarkable features were found \\cite{reviews}. In particular, the higher-order spatial curvature terms can give rise to a bouncing universe \\cite{Calcagni}, may ameliorate the flatness problem \\cite{KK} and lead to caustic avoidance \\cite{Mukohyama:2009tp}. The anisotropic scaling provides a solution to the horizon problem and generation of scale-invariant perturbations with \\cite{inflation} or without inflation \\cite{Mukohyama:2009gg}. It also provides a new mechanism for generation of primordial magnetic seed field \\cite{MMS}, and a modification of the spectrum of gravitational wave background via a peculiar scaling of radiation energy density \\cite{MNTY}. With the projectability condition, the lack of a local Hamiltonian constraint leads to ``dark matter as an integration constant'' \\cite{Mukohyama:2009mz}. The dark sector can also have its purely geometric origins \\cite{Wanga}. A large Non-Gausianity is possible for both scalar \\cite{HW2} and tensor  \\cite{HWY} perturbations even with a single slow roll scalar field, because of the presence of high order derivative terms, and so on.  Despite all of these remarkable features, the theory also faces  some challenging questions, such as instability and strong coupling. To overcome these questions, various models have been proposed, including the ones with an additional local U(1) symmetry \\cite{HMT,zhu}, in which the problems, such as instability, ghosts, strong coupling, and different speeds in the gravitational sector,  can be avoided. In all of those models, the tensor  perturbations are almost the same \\cite{Wang}, so without loss of the generality, we shall work with the model proposed in \\cite{zhu}.  The rest of the paper is organized as follows: In Sec. II we specify the model that accounts for the polarization of PGWs, while in Sec. III we consider the polarization of PGWs in the de Sitter background. In Sec. IV, we discuss their detectability for the Planck satellite and forthcoming observations. The paper is ended with Sec. V, in which we derive our main conclusions.  It should be noted that,  although our motivations  to study the polarization of PGWs is the HL theory, our final conclusions are applicable to any theory in which the dispersion relation of the PGWs is described by Eq.(\\ref{eq14}).  In addition,   the effects of  chirality  of gravitons on  CMB was first studied in \\cite{lue} in Einstein's theory of gravity, and lately   in \\cite{soda} in the HL theory. But, our model is fundamentally different from theirs.   In particular,   the model studied in \\cite{soda}  produces a negligible polarization, and is  not detectable within current and near future observations, as to be shown below.  To have a sizable effect, we find that it is essential for the existence of a non-WKB region in the dispersion relation that leads to non-adiabatic evolution of the modes, a feature that is absent in the model of  \\cite{soda}.   In addition,    chirality  of gravitons  was also considered in  \\cite{Myung}, but this model is not power-counting renormalizable, and cannot be considered as a viable candidate of quantum gravity. ",
        "conclusions": "In this paper,  we have studied the evolution of PGWs, described by the dispersion relation (\\ref{eq14}), obtained from the HL theory of quantum gravity \\cite{zhu}. From the analytical results given by  Eqs.(\\ref{eq71d}), (\\ref{eq71e}) and (\\ref{eq71f}), one can see that the polarization of PGWs is  precisely due to the parity violation and non-adiabatic evolution of the mode function in Region II of Fig. \\ref{fig0}, in which particles are created, where  their occupation number is given by $n_{k} = |\\beta_{k}|^2$. Fig. \\ref{fig3}, on the other hand, shows clearly  that the polarization is considerably enhanced by the third- and four- order spatial derivative terms of Eq.(\\ref{eq14}). The effects of the fifth-order terms were studied in \\cite{soda}, and are showed explicitly in the present paper that their contributions to the polarization are sub-dominant, and  are quite difficult to be detected in the near future. The detectability of the polarization caused by other terms, on the other hand, seems  very optimistic, as shown  in  Fig. \\ref{figure4} (Low panel).  It should be noted that, although the dispersion relation (\\ref{eq14}) is obtained from the HL theory \\cite{zhu,Wang}, our results are actually applicable to any theory where the mode function of PGWs are described by Eqs.(\\ref{eq13}) and (\\ref{eq14}), including the trans-Planck physics \\cite{martin}. In addition,   the  chirality  of PGWs could also be detected by the potential lensing observations \\cite{BKS}."
    },
    "1208/1208.6200_arXiv.txt": {
        "abstract": "The Fermi-LAT telescope has unexpectedly discovered GeV $\\gamma$-ray emission from the symbiotic Nova V407 Cygni. We investigate the radiation processes due to electrons and hadrons accelerated during the explosion of this Nova. \\komm{We consider a scenario in which} GeV $\\gamma$-ray emission observed by Fermi is produced by the electrons with energies of a few tens of GeV in the inverse Compton scattering of stellar radiation. On the other hand, the hadrons are expected to reach larger energies, due to the lack of radiation losses during acceleration process, producing TeV $\\gamma$-rays and neutrinos. We predict the fluxes of very high energy $\\gamma$-rays and neutrinos from Novae of the V407 type for two models of hadron acceleration and discuss their possible detectability by the present and future telescopes (e.g. IceCube, CTA). ",
        "introduction": "Nova V407 Cygni belongs to the class of symbiotic Novae which contains a white dwarf and a red giant (RG) star of the Mira-type. \\komm{The binary period of this system is 43~yrs \\citep{mms90}, and the stars are separated by $\\sim$10$^{14}$ cm.} RG has a radius $\\sim$500 R$_\\odot$, the mass loss rate $3\\times 10^{-7}$ M$_\\odot$ yr$^{-1}$, the wind velocity $\\sim$10 km s$^{-1}$ and the luminosity of $10^4$ L$_\\odot$~\\komm{\\citep{abdo10}}. The matter of the RG wind is accreted onto the surface of a white dwarf. It is commonly accepted that Nova outbursts occur on the surface of a white dwarf as a result of thermonuclear burning of this matter. It is expected that during the explosion a mass of between a few $10^{-7}$ and  a few $10^{-5}$ M$_\\odot$ \\citep{or11} has been expelled with a velocity of $2760-3200$ km s$^{-1}$. This mass interacts with the wind of RG providing conditions for acceleration of particles. It is estimated that such symbiotic Novae may appear within the Galaxy with the frequency of $0.5-5$ yr$^{-1}$\\citep{lu11}. The binary system V407 Cygni is located at a distance of 2.7 kpc \\cite{mms90}  Recently, a $\\gamma$-ray outburst has been observed from the direction of the binary system V407 Cygni by the Fermi-LAT telescope~\\citep{abdo10}. The $\\gamma$-ray emission was detected during the optical outburst, with the differential spectrum well described in the energy range $0.1-6\\,$GeV by a power law with the spectral index $-1.5\\pm0.2$ followed by the exponential cutoff at $E_c=2.2\\pm0.8\\,$GeV. The emission was detected for $\\sim 15$ days (during 2010 March 10-26) with the maximum at 3-4 days after the maximum of the optical emission \\citep{nk10}. Moreover, the $\\gamma$-ray flux at the early stage (10 March - 14 March 2010) and late stage (14 March - 29 March) was different by a factor $\\sim$2 while the spectral features remained unchanged within statistical limits~\\citep{abdo10}. V407 Cyg has been also observed by VERITAS at the later phase of the outburst during 2010 March 19-26~\\citep{aliu12}. Only the upper limit on the flux at the level of $2.3\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ has been reported at the energy of 1.6 TeV.  The detected GeV $\\gamma$-ray emission could have been produced in the Inverse Compton (IC) scattering of the RG soft photons by relativistic electrons or as a result of decay of pions produced in hadronic interactions~\\citep{abdo10}. In fact, modeling of the Novae explosions shows that conditions are convenient for both radiation processes to occur \\citep{or11}. Both radiation processes can provide satisfactory fitting of the $\\gamma$-ray spectrum measured by the Fermi ~\\citep{abdo10}. If GeV $\\gamma$-ray emission is due to the hadronic processes, then neutrinos with $\\sim 10\\,$GeV energies should be also produced as a result of decay of charged pions and muons. It has been argued that such low energy neutrinos (with energies $\\sim0.1-100\\,$GeV) from symbiotic Novae can be detected by the current and future experiments \\citep{raz10}.  In this paper we investigate whether symbiotic Novae can also produce TeV $\\gamma$-rays and neutrinos detectable by the present and future instruments. We assume that the GeV emission observed by Fermi is due to the IC scattering of the RG radiation by relativistic electrons. Their maximum energies are determined by radiation energy losses. However, hadrons might be accelerated to significantly larger energies due to the lack of saturation of the acceleration process by energy losses. These hadrons might produce TeV $\\gamma$-rays and neutrinos in the interaction with the matter of the stellar wind. ",
        "conclusions": "We have considered a two-component scenario for the high energy processes during symbiotic Nova explosions, proposing that GeV $\\gamma$-ray emission is produced by electrons and the accompanying TeV $\\gamma$-ray and neutrino emission is produced by hadrons. \\komm{The GeV $\\gamma$-ray spectral shape seems to favor monoenergetic electron injection, such as can be obtained e.g. in magnetic reconnection process, rather than power-law electron production typical in e.g. Fermi II acceleration.} Two models are considered for the acceleration of the hadrons. In the first model, hadron acceleration process is tuned to the acceleration process of electrons which is constrained by the GeV $\\gamma$-ray observations. This model predicts relatively low fluxes of TeV $\\gamma$-rays which are marginally consistent with the recent upper limits by VERITAS observations provided that similar power goes on the acceleration of electrons and protons. With future CTA project the TeV component in gamma-ray spectrum can be detected even if the luminosity of protons is as low as 3\\% of the one of electrons. The expected neutrino fluxes are below sensitivities of the IceCube telescope. In the second model, hadrons are accelerated independently from electrons in a more efficient mechanism (Fermi I for protons in contrast to Fermi II or reconnection for electrons). This rather optimistic scenario predicts strong fluxes of TeV $\\gamma$-rays which in fact may only appear  during the early stage of the Nova flare (before March 19th) due to the constraints set by the VERITAS upper limits. Also the accompanying neutrino burst at TeV energies might be detectable in this case by the IceCube telescope."
    },
    "1208/1208.1616_arXiv.txt": {
        "abstract": "Multi-field models of inflation predict an inequality between the amplitude $\\tnl$ of the collapsed limit of the four-point correlator of the primordial curvature perturbation and the amplitude $\\fnl$ of the squeezed limit of its three-point correlator. While a convincing detection of non-Gaussianity through the squeezed limit of the three-point correlator would rule out all single-field models, a robust confirmation or disproval of the inequality between $\\tnl$ and $\\fnl$ would provide crucial information about the validity of multi-field models of inflation. In this paper, we discuss to which extent future measurements of the scale-dependence of galaxy bias can test multi-field inflationary scenarios. The strong degeneracy between the effect of a non-vanishing $\\fnl$ and $\\tnl$ on halo bias can be broken by considering multiple tracer populations of the same surveyed volume. If halos down to $10^{13}M_\\odot/h$ are resolved in a survey of volume $25$(Gpc$/h)^3$, then testing multi-field models of inflation at the 3-$\\sigma$ level would require, for instance, a detection of $\\tnl$ at the level of $\\tnl\\sim 10^5$ given a measurement of a local bispectrum with amplitude $\\fnl\\sim 10$. However, we find that disproving multi-field models of inflation with measurements of the non-Gaussian bias only will be very challenging, unless $|\\fnl|\\gtrsim 80$ and one can achieve a halo mass resolution of $\\sim 10^{10}M_\\odot/h$. ",
        "introduction": "Inflation (see \\cite{lrreview} for a review) has  become the dominant  paradigm for understanding the initial conditions for the large scale structure (LSS) formation and for Cosmic Microwave Background anisotropy (CMB). In the inflationary picture, primordial densities are created from quantum fluctuations ``redshifted'' out of the horizon during an early period of superluminal expansion of the universe, where they are ``frozen''. Perturbations at the surface of last scattering are observable as temperature anisotropy in the CMB. The last and most impressive confirmation of the inflationary paradigm has been recently provided by the data of the Wilkinson Microwave Anisotropy Probe (WMAP) mission which has marked the beginning of the precision era of the CMB measurements in space (\\cite{wmap7}).  Despite the simplicity of the inflationary paradigm, the mechanism by which  the cosmological curvature perturbation is generated  is not yet fully established. In the single-field models of inflation, the observed density perturbations are induced by fluctuations of the inflaton field itself. An alternative to the standard scenario is represented by the curvaton mechanism (\\cite{curvaton1}, \\cite{LW}, \\cite{curvaton3})  where the final curvature perturbations are produced from an initial isocurvature perturbation associated to the quantum fluctuations of a light scalar field (other than the inflaton), the curvaton, whose energy density is negligible during inflation. The curvaton isocurvature perturbations are transformed into adiabatic ones when the curvaton decays into radiation much after the end of inflation. Alternatives to the curvaton model are those models characterised by the curvature perturbation being generated by an inhomogeneity in the decay rate (\\cite{rate1}, \\cite{rate2})  of the particles responsible for the reheating after inflation. Other opportunities for generating the curvature perturbation occur at the end of inflation (\\cite{end1}, \\cite{end2}) and during preheating (\\cite{during}). A precise measurement of the spectral index $n_\\zeta$ of the comoving curvature perturbation $\\zeta$ will provide a powerful constraint to single-field models of inflation which predict  the spectral index to be  close to unity. However, alternative mechanisms, like the curvaton, also predict a value of the spectral index  very close to unity. Thus, even a precise measurement of the spectral index will not allow us to efficiently distinguish among them. Furthermore, the lack of a gravity-wave signal in CMB anisotropies would not give us any information about the perturbation generation mechanism, since alternative mechanisms predict an amplitude of gravity waves far too small to be detectable by future experiments aimed at observing the $B$-mode of the CMB polarisation.  There is, however, a third observable which will prove fundamental in providing information about the mechanism chosen by Nature to produce the structures we see today. It is the deviation from a Gaussian statistics, {\\it i.e.}, the presence of higher-order connected correlation functions of the perturbations. Indeed, a possible source of non-Gaussianity (NG) could be primordial in origin, being specific to a particular mechanism for the generation of the cosmological perturbations (for a review see \\cite{BKMR}). This is what makes a positive detection of NG so relevant: it might help discriminating among competing scenarios which, otherwise, would might remain indistinguishable.  To characterise the level of NG in the comoving curvature perturbation, one usually introduces two nonlinear parameters,  $\\fnl$ and $\\tnl$. The first one is defined in terms of the three-point correlator, the bispectrum, of the comoving curvature perturbation in the so-called squeezed limit \\be \\fnl=\\frac{5}{12}\\frac{\\langle \\zeta_{\\vk_1}\\zeta_{\\vk_2}\\zeta_{\\vk_3}\\rangle^\\prime} {P^\\zeta_{\\vk_1}P^\\zeta_{\\vk_2}}\\,\\,\\,\\,\\,\\,\\,\\,(k_1\\ll k_2\\sim k_3)~. \\ee The second one is defined  in terms of the four-point correlator, the trispectrum, in the so-called collapsed limit \\be \\tnl=\\frac{1}{4}\\frac{\\langle \\zeta_{\\vk_1}\\zeta_{\\vk_2}\\zeta_{\\vk_3}\\zeta_{\\vk_4}\\rangle^\\prime} {P^\\zeta_{\\vk_1}P^\\zeta_{\\vk_3}P^\\zeta_{\\vk_{12}}}\\,\\,\\,\\,\\,\\,\\,\\,(\\vk_{12}\\simeq  0)~. \\label{tf2} \\ee We have normalised the correlators with respect to the power spectrum of the curvature perturbation, \\be \\langle\\zeta_{\\vec{k}_{1}}\\zeta_{\\vec{k}_{2}}\\rangle = (2\\pi)^3\\delta({\\vec{k}_{1}}+{\\vec{k}_{2}})P^\\zeta_{\\vec{k}_{1}} \\ee and used the notation  $\\vec{k}_{ij}=(\\vec{k}_{i}+\\vec{k}_{j})$. In all single-field models of inflation the bispectrum is suppressed in the squeezed limit and is non vanishing only when the spectral index deviates from unity, $\\fnl=5/12(1-n_\\zeta)\\simeq 0.02$ (see \\cite{acqua}, \\cite{con1},\\cite{con2}, \\cite{con3}). A convincing detection of NG in the squeezed limit, $\\fnl\\gg 1$, would therefore rule out all single-field models (one should be aware though that, in single-field models of inflation, a large NG can be generated in shapes others than the squeezed, {\\it e.g.} in the equilateral configuration). However, such a detection would not rule out multi-field models of inflation where the NG is seeded by light fields other than the inflaton. How can  we derive some useful informations about them? In this respect, the collapsed limit of the four-point correlator is particularly important because, together with the squeezed limit of the three-point correlator, it may lead to the so-called Suyama-Yamaguchi (SY) inequality (\\cite{SY}, see also \\cite{SY1}, \\cite{SY2}). Based on  the  conditions  that 1) scalar fields are responsible for generating curvature perturbations and  that 2)  the fluctuations in the  scalar fields at the horizon crossing are scale invariant and Gaussian, Suyama and Yamaguchi proved the inequality \\be \\label{SYin} \\tnl\\ge \\left(\\frac{6}{5}\\fnl\\right)^2. \\ee The  condition 2) amounts to assuming that the connected three- and four-point correlations of the light  fields vanish and that the NG is generated at super-horizon scales. This is quite a restrictive assumption. However, based on the operator product expansion, which is particularly powerful in characterising in their full generality the squeezed limit of the three-point correlator and the collapsed limit of the four-point correlator, it was shown that the SY inequality holds also for NG light fields (\\cite{KR}). This  is consequence of fundamental physical principles (like positivity of the two-point function) and its hard violation would require some new non-trivial physics to be involved.  The observation of a strong  violation of the inequality will then have profound implications for inflationary models. It will imply either that  multi-field inflation cannot be responsible for generating the observed fluctuations independently of the details of the model, or that some new non-trivial (ghost-like) degrees of freedom play a role during inflation (\\cite{KR}).  Testing the SY inequality with future LSS observations and, therefore, the validity of multi-field inflationary models is the subject of this paper. The squeezed limit of the bispectrum and the collapsed limit of the trispectrum are  particularly interesting from the observationally point of view because they are associated to pronounced effects of NG on the clustering of dark matter halos and, in particular, to a strongly scale-dependent bias (\\cite{dalal}). Measurements of the galaxy power spectrum have been exploited to set limits on primordial non-Gaussianity competitive with those inferred from CMB observations (\\cite{slosar}, \\cite{DS10}, \\cite{X11}). As we have seen, a large value of $\\fnl$ in the squeezed limit implies that the cosmological perturbations are generated within a multi-field model of inflation where the NG is sourced by light fields other than the inflaton. An inescapable consequence of the SY inequality (\\ref{SYin}) is that the NG is also characterised by a large trispectrum in the collapsed limit. Therefore, investigations that take advantage of the scale-dependent effects of NG on the clustering of dark matter halos should in principle take into account both $\\fnl$ and $\\tnl$. However, since the contribution from the latter is suppressed by $10^{-4}(\\tnl/\\fnl)$, setting limits on $\\fnl$ under the assumption $\\tnl=0$, as done in the literature, should be a good approximation unless $|\\tnl|\\gg \\fnl^2$.  In this paper, we will essentially try to answer the following question: what values of $\\fnl$ and $\\tnl$ have to be measured in order to either confirm or disprove the SY inequality ? As we shall see, even though the contributions from $\\fnl$ and $\\tnl$ are degenerate in the non-Gaussian halo bias, combining multiple halo mass bins can greatly help breaking the degeneracy. As we shall demonstrate, testing  multi-field models of inflation at the 3-$\\sigma$ level would require, for a {\\small EUCLID}-like survey, a detection of a four-point correlator amplitude in the collapsed limit of the order of $\\tnl\\sim 10^5$ given a measurement of a local bispectrum at the level of $\\fnl\\sim 10$. Conversely, we will argue that disproving multi-field models of inflation would require a detection of $|\\fnl|$ at the level of 80 or larger if dark matter halos can be resolved down to a mass $10^{10}M_\\odot/h$.  The paper is organised as follows. Section 2 contains a short summary of the impact of primordial NG on the halo bias at large scales. Section 3 describes the methodology adopted. The last Section presents the results and discusses their implications. In all illustrations, the cosmology is a flat $\\Lambda$CDM Universe with normalisation $\\sigma_8=0.807$, hubble constant $h_0=0.701$ and matter content $\\Omega_{\\rm m}=0.279$. ",
        "conclusions": "\\label{sec:results}  We first compute the uncertainties on $\\fnl$ and $\\tnl$ from two different tracer populations and for a shot-noise matrix consistent with Poisson noise, i.e. ${\\bf E}={\\rm diag}(1/n(M_1),1/n(M_2))$. We consider a nearly unbiased sample with average mass $M\\sim 10^{12}M_\\odot/h$ and a high mass sample with $M=10^{14}M_\\odot/h$. Table \\ref{tab:data} summaries the characteristics of these populations. For a given mass $M$, the second-order Lagrangian bias parameter $b_2(M)$ is computed from the Sheth-Tormen multiplicity function, whereas the skewness $S_{s,{\\rm loc}}^{(3)}$ is computed from the phenomenological relation given in \\S\\ref{sec:ngbias}. Fig.\\ref{fig:ellipse} shows the resulting 68, 95 and 99\\% confidence contours for the parameters $\\fnl$ and $\\tnl$ when the fiducial model assumes $\\fnl=\\pm 10$ and $\\tnl = 2\\times 10^4$. The 1-$\\sigma$ errors are $\\sigma_{\\fnl}\\simeq 23$ and $\\sigma_{\\tnl}\\simeq 2.0\\times 10^5$. We have tried different combinations of halo populations and found that the errors do not change significantly. At this point, we would conclude that galaxy bias alone cannot yield interesting constraints on $\\tnl$ and $\\fnl$.  The situation changes dramatically when the surveyed halos are divided into $N\\gg 1$ populations of increasing mass, with equal number density. In Fig. \\ref{fig:mmin}, symbols represent the halo model prediction for the 1-$\\sigma$ uncertainties $\\sigma_{\\fnl}$ and $\\sigma_{\\tnl}$ in the limit of infinitely many halo bins. The shot-noise matrix now takes the form Eq.(\\ref{eq:snmatrix}). Red triangles indicate $\\sigma_{\\fnl}$ in a one-parameter model with $\\fnl=0$ (left panel) and $\\fnl=10$ (right panel). Filled and empty squares represent $\\sigma_{\\fnl}$ and $\\sigma_{\\tnl}$ in a two-parameters model with $(\\fnl,\\tnl)=(0,0)$ (left panel) and $(\\fnl,\\tnl)=(10,2\\times 10^4)$ (right panel). Results are shown as a function of the mass of the smallest halos resolved in the survey. Compared to the previous configuration, significant gains are already achieved for $M_{\\rm min}\\approx 10^{13}M_\\odot/h$. While the constraint on $\\fnl$ is somewhat degraded if one allows for a non-zero $\\tnl$, the 1-$\\sigma$ uncertainty on $\\tnl$ is of the order of $(10^3 - 10^4)$, an order of magnitude better than in the case of two galaxy populations. Table \\ref{table:table2} gives the 1-$\\sigma$ errors for $M_{\\rm min}=10^{13}$ and $10^{10}M_\\odot/h$.  How well can we test the SY inequality with galaxy bias? Fig. \\ref{fig:set1} displays, as a function of $\\fnl$,  the minimum value of $\\tnl$ for which the difference $(\\tnl-(36/25)\\fnl^2)$ is greater than its 1-,2- and 3-$\\sigma$ error which, for Gaussian-distributed $\\fnl$ and $\\tnl$, reads \\begin{align} \\sigma^2_{\\tnl-\\frac{36}{25}\\fnl^2} &= \\sigma^2_{\\tnl}+2\\left(\\frac{36}{25}\\right)^2\\sigma_{\\fnl,\\tnl}^2 \\\\ & \\quad + 4\\left(\\frac{36}{25}\\right)^2\\sigma^2_{\\fnl}\\bar{f}^2_{\\rm NL} -4\\left(\\frac{36}{25}\\right)\\sigma_{\\fnl,\\tnl}\\bar{f}_{\\rm NL} \\nonumber \\;, \\end{align} where $\\sigma^2_{\\fnl}, \\sigma_{\\fnl,\\tnl}$ and $\\sigma^2_{\\tnl}$ are the entries of the inverted Fisher matrix and $\\bar{f}_{\\rm NL}$, $\\bar{\\tau}_{\\rm NL}$ are the values of the fiducial model assumed. The various curves indicate the halo model prediction for $N\\gg 1$ halo populations with a minimum resolved mass $M_{\\rm min}=10^{13}M_\\odot/h$. For instance, if a non-vanishing value of $\\fnl=10$ is measured in the future, then the contribution induced by the collapsed limit of the trispectrum must be detected with an amplitude of at least $\\tnl \\sim {\\cal O}(1)\\times 10^5$ in order to have a 3-$\\sigma$ detection of the SY inequality with the non-Gaussian halo bias. Of course, these values are only indicative since the analysis is performed with the restrictive assumption of Gaussian errors.  \\begin{figure} \\center \\resizebox{0.40\\textwidth}{!}{\\includegraphics{fig3.eps}} \\caption{Testing the validity of the SY inequality with measurements of the non-Gaussian bias. The various curves are halo model predictions for $N\\gg 1$ halo mass bins with $M_{\\rm min}=10^{13}M_\\odot/h$. At fixed value of $\\fnl$, they indicate the minimum $\\tnl$ required in order to have a measurement of the SY inequality at the 1-, 2- and 3-$\\sigma$ confidence level.} \\label{fig:set1} \\end{figure}  Finally, we can also assess how well galaxy bias can probe the violation of the SY inequality. As stated above, the observation of a strong  violation would have profound implications for inflationary models as it implies either that multi-field inflation, independently of the details of the model, cannot be responsible for generating the observed fluctuations, or that some new non-trivial (ghost-like) degrees of freedom play a role during inflation. Measuring a violation essentially consists in a simultaneous detection of  a non-zero value of $\\fnl$ and a (non-zero) small enough value of $\\tnl$. Here, we have simply estimated the smallest $|\\fnl|$ such that $(36/25)\\fnl^2$ is larger than the 3-$\\sigma$ error on $\\tnl$. Having found that, for the current observationally allowed range of $\\fnl$, the error of $\\tnl$ does not significantly change if we set in all runs $\\tnl=0$, we have thus computed $\\sigma_{\\tnl}$ assuming a vanishing value of $\\tnl$. A comparison of $3\\sigma_{\\tnl}$ with $(36/25)\\fnl^2$ shows that, for a minimum halo mass $M_{\\rm min}=10^{13}M_\\odot/h$, the SY inequality cannot be tested with the non-Gaussian galaxy bias solely for realistic values of $\\fnl$. Even if halos are resolved down to $10^{10}M_\\odot/h$ is $3\\sigma_{\\tnl} < (36/25)\\fnl^2$ satisfied only for $|\\fnl|$ larger than $\\sim$ 80.  Summarising,  a large NG in the squeezed limit implies that the cosmological perturbations are generated by some light scalar field other than the inflaton. The  SY inequality (\\ref{SYin}) inevitably imposes that a large trispectrum in the collapsed limit is also present. However, the contribution of $\\tnl$ to the non-Gaussian halo bias is suppressed by $10^{-4}(\\tnl/\\fnl)$ and strongly degenerate with that induced by $\\fnl$. Notwithstanding this, we have shown that multi-tracer methods can exploit the distinct mass-dependence of the $\\fnl$- and $\\tnl$-induced bias corrections to reduce the 1-$\\sigma$ uncertainty down to $\\sigma_{\\tnl}\\lesssim 10^4$ (and simultaneously achieve $\\sigma_{fnl}\\sim 1-5$) for a survey covering half of the sky up to $z\\approx 1$. The exact values depend on the mass $M_{\\rm min}$ of the least massive halos observed. Our results on the capability of testing  the SY inequality through the NG scale-dependent bias are summarised in Fig.\\ref{fig:set1}. The latter shows that testing the SY inequality at the level of 3-$\\sigma$ would require detecting $\\tnl$ at the level of $10^5$ for the minimum resolved mass  $M_{\\rm min}=10^{13}M_\\odot/h$. Conversely, testing the violation of the SY inequality requires both a much smaller resolved mass, $M_{\\rm min}=10^{10}M_\\odot/h$ and a large bispectrum, $|\\fnl|\\gtrsim 80$. As mentioned above, all these results are valid provided that $\\gnl=0$ and that the nonlinear parameters $\\fnl$ and $\\tnl$ estimated from the data are Gaussian-distributed. Relaxing these assumptions will be the subject of future work."
    },
    "1208/1208.3425_arXiv.txt": {
        "abstract": "We present analysis of archival X-ray data obtained with the \\xmm and \\suzaku for a new Intermediate Polar identified as a counterpart of an INTEGRAL discovered $\\gamma$-ray source, \\emph{IGR J17195-4100}.  We report a new period of 1053.7$\\pm$12.2~s in X-rays. A new binary orbital period of 3.52$^{+1.43}_{-0.80}$~h is strongly indicated in the power spectrum of the time series. An ephemeris of the new period proposed as the spin period of the system has also been obtained. The various peaks detected in the power spectrum suggest a probable disc-less accretion system. The soft X-rays ($<$3~keV) dominate the variability seen in the X-ray light curves.  The spin modulation shows energy dependence suggesting the possibility of a variable partial covering accretion column. The averaged spectral data obtained with  \\xmm EPIC cameras show a multi temperature spectra with a soft excess. The latter can be attributed to the varying coverage of accretion curtains. ",
        "introduction": "Cataclysmic variables (CVs) are interacting binaries with a white dwarf accreting matter from a main sequence companion star through a Roche lobe, with a typical orbital period of few hours (see \\Citealp{warner95} for a review).  Many of the CVs have high magnetic fields and are called as magnetic CVs (mCVs). The mCVs are further classified as polars \\Citep[see][]{cropper90}, and Intermediate polars (IPs; see \\Citealt{patterson94},\\Citealt{hellier96}) based on the strength of the observed magnetic fields in them. Polars have very high magnetic fields (10-80~MG) which do not allow the formation of accretion discs, but help to synchronise the spin and orbital periods to a very high degree. Intermediate polars have somewhat  less magnetic fields compared to polars ($\\sim$ 10 MG) but sufficient to disrupt the accretion disc and  truncate it at a certain distance from the white dwarf. The IPs typically have a spin period of few hundreds of seconds and orbital periods of few hours and many IPs cluster around a P$_{spin}$/P$_{orb}$ $\\approx$ 0.05 - 0.15 \\Citep{norton04}. The accreting matter in both polars and IPs falls on to the white dwarf through poles along the magnetic field lines. A few IPs show disc-less accretion, and in some cases accretion stream is seen to flip between the two poles as was first reported in V2400 Oph by \\Citet{buckley97}. Generally the IPs are hard X-ray emitters.  However, many IPs also show a soft X-ray excess typically of a blackbody component in soft X-rays (\\Citealt{evans07}, \\Citealt{martino08}, \\Citealt{anzolin09}). The soft blackbody component in IPs is usually attributed to the presence of accretion curtains \\Citep{evans07}.   The INTEGRAL/IBIS soft Gamma-ray  survey \\Citep{barlow06,bird07} has detected a total of 32 CVs till now which is close to 10 per cent of INTEGRAL detections. These detections are identified  mainly with mCVs with 22 confirmed or probable IPs and three polars. IGR J17195-4100 (IGR1719 hereafter) was detected  as a source in  \\integral observations by \\Citet{bird04} and was later identified with an optical object of R$_{mag}$ $\\approx$ 14.3~mag and classified as an IP  by \\Citet{masetti06} based on its optical spectrum.  \\Citet{butters08} reported candidate periods of 1842.4$\\pm$1.5~s and 2645.0$\\pm$4.0~s in \\igr based on RXTE observations, whereas a spin period of 18.9925$\\pm$0.0006~min (1139.55$\\pm$0.04~s) and an orbital period of 4.005$\\pm$0.006~h were reported by \\Citet{pretorius09} from high speed optical photometry. \\citet{masetti06} determined the distance of \\igr close to 110~pc assuming an absolute magnitude M$_V$ $\\sim$ 9. The X-ray luminosity of the IP as given by \\citet{masetti06} is 3.6$\\times$10$^{31}$~erg~s$^{-1}$ in the 0.5-10.0~keV energy range and 5.5$\\times$10$^{31}$ erg~s$^{-1}$ in the 20-100~keV energy range. \\Citet{yuasa10} analyzed wide-band \\emph{Suzaku} spectra of seventeen IPs with \\igr being one of them.  The mass, shock temperature and spectral parameters were derived by fitting the X-ray spectra with numerically calculated emission models and the mass of the white dwarf in \\igr was estimated as 1.03$^{+0.24}_{-0.22}$~M$_\\odot$ \\citep{yuasa10}. ",
        "conclusions": "Our timing analysis of the archival \\xmm and \\suzaku data of IGR1719 shows a dominant frequency at 0.949$\\pm$0.011~mHz corresponding to a period of 1053.7$\\pm$12.2~s and another close period at 1149.4~s. We propose that the period at 1053.7~s is the spin period of the white dwarf, and the period at 1149.4~s is the orbital sideband, and thus we estimate the orbital period as 3.52$^{+1.43}_{-0.80}$~h. We also provide an  ephemeris of this spin period.  Though the results of the phase resolved spectroscopy performed at the two periods 1053.7~s and 1149.4~s favours the proposed period at 1053.7~s as the white dwarf spin period, polarimetric detection of the spin period of \\igr is needed to confirm the spin period. The power spectra obtained in different energy bands suggest that \\igr is a candidate for the rare group of disc-less IPs.  A partially and fully absorbed cooling flow model satisfactorily explains the observed spectra of IGR1719.  The partial covering fraction varies between 0.34 at the spin minima to 0.46 at the spin maxima indicating a geometrical change in the absorber with spin phase. We also report a soft X-ray excess which can be modelled with a low kT \\apec model.  We detect a 6.4~keV fluorescent line which can be attributed to the emission from cold iron in the ionization states up to Fe~XVII from the white dwarf surface. From the observed flux in the 0.3-10.0~keV energy range, we estimate a mass accretion rate of 5.5$\\times$10$^{-12} M_\\sun$~yr$^{-1}$.  From the temporal and spectral analysis the image of \\igr that emerges is of a disc-less accreting system where the matter from the secondary  is threaded along a magnetic pole as a stream resulting in modulations of X-rays at the orbital beat period. The shocks formed above the white dwarf surface give a multi temperature spectra with a shock temperature above 64~keV. The soft X-ray excess seen in \\igr may be emitted by the unhindered accretion footprint. The site for origin of the fluorescent line emission is not clear in the system.  ~\\\\~  \\noindent\\emph{Acknowledgements:} This research has made use of data obtained from the High Energy Astrophysics Science Archive Research Center (HEASARC), provided by NASA's Goddard Space Flight Center.  We would like to thank the anonymous referee for the insightful remarks that improved the manuscript."
    },
    "1208/1208.3431_arXiv.txt": {
        "abstract": "% We report the first analyses of SDO/AIA observations of  the formation of a quiescent polar crown prominence in a coronal cavity. The \\ion{He}{2} 304~\\AA\\ ($\\log\\mathrm{T_{max}}\\sim 4.8$~K) data show both the gradual disappearance of the prominence due to vertical drainage and lateral transport of plasma followed by the formation of a new prominence 12 hours later. The formation is preceded by the appearance of a bright emission ``cloud'' in the central region of the coronal cavity. The  peak brightness of the cloud progressively shifts in time from the \\ion{Fe}{14} 211~\\AA\\  channel, through the \\ion{Fe}{12} 193~\\AA\\  channel, to the \\ion{Fe}{9} 171~\\AA\\  channel ($\\log\\mathrm{T_{max}}\\sim 6.2, 6.1, 5.8$~K, respectively) while simultaneously decreasing in altitude. Filter ratio analysis estimates the initial temperature of the cloud to be approximately $\\log\\mathrm{T}\\sim 6.25$~K with evidence of cooling over time. The subsequent growth of the prominence is accompanied by darkening of the cavity in the 211~\\AA\\ channel. The observations imply prominence formation via \\emph{in situ} condensation of hot plasma from the coronal cavity, in support of our previously proposed process of magneto-thermal convection in coronal magnetic flux ropes. ",
        "introduction": "\\label{sect_intro}   Quiescent prominences, or filaments, are dynamic formations of relatively cool and dense plasma (T $\\sim 10^4$~K, $\\mathrm{n_{e}}\\sim 10^{11}\\pcmc$) extending into the much hotter and rarefied corona (T $\\sim 10^6$~K, $\\mathrm{n_{e}}\\sim 10^{8}\\pcmc$) over magnetic polarity boundaries far from active regions \\citep{Martin:1998,TandbergHanssen:1995ud,Priest:1988,Schmieder:1984ul}. In the ``polar crown'' regions they are frequently observed in so-called ``coronal cavities,'' dark elliptical structures seen during eclipses or in extreme ultraviolet (EUV) and X-ray images of the corona \\citep{Kucera:2012,Reeves:2012bu,Schmit:2011ix,Habbal:2010,Fuller:2009jm,Gibson:2006kr}.  Coronal cavities and prominences are the progenitors of coronal mass ejections (CMEs), with the overlying arcade loop system, the cavity, and the prominence comprising the leading dense shell, dark void, and bright trailing material, respectively \\citep{Dere:1999ki,Gibson:1998ua,Illing:1986gq}.  Understanding the mechanisms of formation, evolution, and eventual eruption of this complex magnetic system is a central goal of solar physics.  Prominence formation mechanisms can be generally classified into chromospheric transport and coronal condensation \\citep{Labrosse:2010bt,Mackay:2010fp}.  In the former, \\emph{cool} plasma is lifted from the chromosphere into the corona by siphon flows \\citep{PikelNer:1971ft}, magnetic reconnection \\citep{Litvinenko:1999gc}, or magnetic flux emergence \\citep{Okamoto:2008kg}.  In the latter, \\emph{hot} plasma \t% in the local corona condenses into a prominence \\citep{Xia:2012,Luna:2012eb,Karpen:2008hc,Antiochos:1999ef,An:1985fa,1983SoPh...88..219P}. Confirming observations of any of these mechanisms have thus far been elusive.  Here we report the first analysis of {\\it Solar Dynamics Observatory} Atmospheric Imaging Assembly \\citep[SDO/AIA;][]{Lemen:2011} observations of  apparent \\emph{in situ} formation of a quiescent prominence within a coronal cavity. The results are consistent with earlier studies of cloud prominence formation \\citep{Liu:2012gq} and support the hypothesis that coronal cavities and prominences are not in magnetostatic equilibrium, but instead harbor a novel magneto-thermal convection \\citep{Berger:2011eoa} with hot upflows and condensation downflows enabling the build-up of magnetic energy and helicity and culminating in an eruptive CME \\citep{Zhang:2005dn}.  Recent theoretical analyses \\citep{Low:2012tu,Low:2012tv} suggest that MHD condensation has a resistive origin due to runaway thermal collapse of condensing plasmas and spontaneous current sheets \\citep{Parker:1994cs,Parker:1953}, leading to rapid condensation and cross-field transport of prominence plasma. ",
        "conclusions": "% \\label{sect_discussion}  We have analyzed observations of the dynamic formation of a quiescent polar crown prominence in a coronal cavity. The sequence of events is summarized as follows: \\begin{enumerate} \\item A pre-existing prominence slowly disappears due mostly to drainage and lateral transport of plasma (Fig.~\\ref{fig:171+304}). \\item A bright emission cloud forms in the upper regions of the coronal cavity (Figs.~\\ref{fig:171+304} and \\ref{fig:193+171}). \\item The cloud descends toward the lower region of the cavity while sequentially becoming brighter in the 211, 193, and 171~\\AA\\ channels (Fig.~\\ref{fig:timeslice}). \\item A new prominence appears in the 304~\\AA\\ channel and rapidly grows in both vertical and horizontal extent (Figs.~\\ref{fig:171+304} and \\ref{fig:timeslice}). \\item The coronal cavity core above the prominence grows darker in the 211~\\AA\\ channel as the 304~\\AA\\ prominence grows (Fig.~\\ref{fig:mass}). \\item When the prominence reaches its maximum size after approximately 18~h of growth, the emission cloud in the cavity is completely gone (Fig.~\\ref{fig:171+304}). \\end{enumerate}  Taken together, these observations are consistent with the hypothesis that the quiescent prominence has formed via condensation from hotter plasma contained in the core of the coronal cavity. We interpret the EUV emission sequences in Figs.~\\ref{fig:171+304},  \\ref{fig:193+171}, and \\ref{fig:timeslice} as evidence of radiative cooling and descent of the hot plasma in the coronal cavity core, first appearing bright in the hotter  211~\\AA\\ channel with a peak temperature of formation of $\\log\\mathrm{T_{max}}\\sim 6.2$~K, followed by a shift to emission in the 193~\\AA\\ channel ($\\log\\mathrm{T_{max}}\\sim 6.1$~K), followed by another shift to emission in the  171~\\AA\\ channel ($\\log\\mathrm{T_{max}}\\sim 5.8$~K), and finally appearing in the ``chromospheric'' 304~\\AA\\ channel ($\\log\\mathrm{T_{max}}\\sim 4.8$~K).  The profile of the hotter filter ratio histogram widens toward cooler temperatures over time (Fig.~\\ref{fig:twofilter}(g)), further supporting our interpretation that the observed plasma radiatively cools to condense into the prominence.  The drop in height of the cloud further supports this hypothesis. Radiative cooling is proportional to $\\mathrm{n_{e}}\\!^2$ at corona temperatures and increases with decreasing temperature in the $\\log\\mathrm{T} \\sim$ 5--7~K range. Therefore a ``runaway thermal instability'' is possible in which the plasma cools increasingly rapidly as it becomes denser and cooler \\citep{Low:2012tu}. The increased density of the cloud naturally results in higher gravitational force countering the local Lorentz force of the cavity magnetic field. Thus the descent of the cloud is consistent with radiative cooling and condensation in a weakly magnetized environment as the plasma seeks a new equilibrium height, possibly through resistive reconnection \\citep{Low:2012tv}. As the prominence gains mass at the expense of the hot cloud, the coronal cavity simultaneously becomes more sharply defined in the 211~\\AA\\ channel due to mass loss from the core region.  The observations shown here imply the possibility of \\emph{in situ} formation of prominences in coronal cavities. Given the large density disparity, \\cite{TandbergHanssen:1995ud} estimated that the entire coronal mass could only support such condensation for a few large prominences.  However this concern was based on the implicit assumption that prominences are magnetostatic suspensions of a fixed amount of plasma. But as shown by recent SDO/AIA and \\emph{Hinode}/SOT observations, quiescent prominences are in constant motion via drainage downflows \\citep{Liu:2012gq,Haerendel:2011de,Chae:2008ca} and re-supply of mass through various mechanisms \\citep[e.g.,][]{Berger:2011eoa,Li:2012,Su:2012}.  \\cite{Liu:2012gq} estimated the mass loss in prominence downflows  to be $10^{15}$~g/day or roughly the equivalent of a typical CME mass. Thus the mass in a prominence changes significantly with time and is not directly related to the mass in the global corona in a simple manner. Mass balance in prominences appears to be a cyclic process in which the mass transported into the coronal cavity determines how much can eventually condense to form prominences. Similar cyclic processes have been proposed for the quiet Sun coronal mass balance \\citep{McIntosh:2012,Antolin:2012,Marsch:2008} and coronal rain condensation \\citep{Landi:2009,Schrijver:2001}.  Observations at the limb necessarily integrate along heliographic longitude. Thus it is possible that the prominence that disappears is not at the same longitude as the new prominence that forms, or that the cloud and the new prominence are a line-of-sight coincidence. But the fact that the shrinkage of the cloud (darkening of the cavity), the descent of the cloud, the formation of prominence ``horns'', and the eventual formation of a PCTR sheath are all synchronized with the growth of the prominence in the 304~\\AA\\ channel argue against a line-of-sight coincidence. We note that the hot cloud itself exhibits no significant flows, other than its slow radial descent, during the condensation process. Thus it is unlikely that we are confusing \\emph{in situ} condensation with flow-based formation mechanisms such as flows between disparate magnetic footpoints \\citep[e.g.,][]{Xia:2012,Luna:2012eb,Karpen:2008hc,Antiochos:1999ef}.              {\\scriptsize"
    },
    "1208/1208.0040.txt": {
        "abstract": "We present measurements of the linear diameter of the emission region of the Vela pulsar at observing wavelength $\\lambda=18\\ {\\rm cm}$.  We infer the diameter as a function of pulse phase from the distribution of visibility on the Mopra-Tidbinbilla baseline. As we demonstrate, in the presence of strong scintillation, finite size of the emission region produces a characteristic W-shaped signature in the projection of the visibility distribution onto the real axis. This modification involves heightened probability density near the mean amplitude, decreased probability to either side, and a return to the zero-size distribution beyond. We observe this signature with high statistical significance, as compared with the best-fitting zero-size model, in many regions of pulse phase. We find that the equivalent full width at half maximum of the pulsar's emission region decreases from more than 400\\ km early in the pulse to near zero at the peak of the pulse, and then increases again to approximately 800\\ km near the trailing edge. We discuss possible systematic effects, and compare our work with previous results. ",
        "introduction": "Pulsars emit strong radio emission from compact regions. The enormous magnetic fields and rapid rotation of neutron stars easily accelerate electrons and positrons to high energy, but the means by which a small fraction of that energy is transformed to radio emission remains poorly understood. Pulsar emission regions are small, but interstellar scattering of radio waves provides an astronomical-unit-scale lens with the nanoarcecond resolution sufficient to resolve them spatially. However, the lens is highly corrupt, and unraveling source structure involves application of statistical models to large volumes of high-quality data. These models must include accurate descriptions of the effects of scattering and of noise.  In this paper, we describe observations and data analysis to fit a simple model for a spatially-extended emission region to the interferometric visibility statistics of the Vela pulsar. We focus on description of the models, the data, and comparisons of the two. In this introductory section, we briefly describe, as background, the features of interstellar scattering important for our technique and the basic features of pulsar physics involved in the interpretation.  \\subsection{Pulsar Emission Physics}  Although pulsars have been observed for more than 40 years, the process by which a rapidly-rotating, magnetized neutron star converts a small fraction of its rotational energy to radio waves remains unclear. The rapid rotation would produce $\\vec v\\times \\vec B$ forces and induced electric fields ample to tear electrons from the surface of the neutron star, were those forces not cancelled by an induced co-rotating charge distribution \\citep{gol69}. The magnetic field of the neutron star away from its surface is nearly dipole, modified by inertia of the corotating particles and relativistic effects \\citep{spi06}. ``Open'' field lines pass through the light cylinder, the surface where the co-rotation speed is that of light. These field lines carry highly-relativistic charged particles away from the star, forming a powerful wind. A small fraction of the energy of this wind is apparently converted to radio emission, observed as pulses because of stellar rotation.  The boundaries of the set of open field lines form promising places for the origin of pulsar emission. Above the ``polar cap'' at the base of the open field lines, the current may be insufficient to replace the outflowing charge, so that a gap may form, with strong electric field parallel to the nearly-vertical magnetic field \\citep{rud75,aro79}. Similarly, a gap may form between the last closed field lines, which nearly graze the light cylinder, and the open field lines. Proposed locations include a ``slot gap\" extending to high altitude from the polar cap \\citep{mus04} and charge-free ``outer gaps\" in the outer magnetosphere \\citep{che86,che00}. Within these gaps, electrons and positrons accelerate to TeV energies, accompanied by pair creation. These particles can emit X-rays and gamma-rays via curvature, synchrotron, and inverse Compton emission. Force-free simulations of pulsar magnetospheres cannot include gaps, but interestingly predict strong currents in the same locations, which might produce particle acceleration via plasma instabilities \\citep{spi06,Gru07}. These particle-acceleration mechanisms may provide the power source for radio emission.  The radio emission is coherent: the observed $\\sim 10^{26}$\\ K brightness temperature exceeds that possible for individual electrons \\citep{Man77}. The nanosecond variability of giant pulses indicates that the emission originates in structures $\\sim 1$ m across \\citep{Han03}. These structures are likely distributed over scales comparable to those of the particle-acceleration zones: larger than the polar cap, $\\sim 1$ km, but smaller than the diameter of the light-cylinder. In this work, we seek to determine the lateral dimension of this region of emission.  Emission could arise directly from particles traveling along the field lines, in which case polarization and temporal variations would directly reflect conditions at the source. Alternatively, emission may be reprocessed before it leaves the pulsar's magnetosphere, as is suggested by observations and theoretical models \\citep{lyu00,Jes10}. Radio emission could also arise via production of plasma waves from plasma instabilities on open field lines. These plasma waves would propagate nearly along field lines, and then be converted to radio waves where the local plasma frequency falls below the observed frequency \\citep{Bar86}. Our measurements can contribute to this picture by describing the lateral scale of the emission region.  \\subsection{Interstellar Scattering and Scintillation}  Radio waves emitted from a pulsar encounter variations in refractive index in the interstellar medium, from variations in electron density. Waves travel onward with ``crinkled'' phase fronts and arrive at an observer from along a number of paths to form a diffraction pattern. For the interstellar medium at decimeter or longer wavelengths, the differences in path lengths are many wavelengths, and many paths contribute to the diffraction pattern at the observer \\citep{coh74}.  The diffraction pattern at the observer is the convolution of an image of the source with the pattern for a point source \\citep{cor89,goo96}. This results from the fact that, for small deflections, the Kirchoff integrals that relate the electric field at the observer to those at the source become Fourier transforms, with the effects of the scattering medium inserted multiplicatively at the screen \\citep{Gwi98}. Geometrical factors, depending on the position of the screen, relate the original source and corrupt image by a magnification factor $M=D/R$, where $D$ is the distance from observer to scatterer, and $R$ is the distance from scatterer to source.  Because path differences are thousands of wavelengths, paths that reinforce at one observing frequency and position in the observer plane may cancel at nearby frequencies or positions. For an observed angular extent $\\theta_{\\rm ISS}$ of the scattering disk and observing wavelength $\\lambda$, the scale of the diffraction pattern at the observer is $S_{\\rm ISS} = \\lambda/\\theta_{\\rm ISS}$, equal to the linear resolution of the scattering disk viewed as a lens. The source shows scintillations, or intensity variations, with timescale $\\Delta t_{\\rm ISS} = S_{\\rm ISS}/V_{\\perp}$ as the line of sight sweeps through scattering material at speed $V_{\\perp}$. A sharp pulse at the pulsar will arrive over a range of times $\\tau_{\\rm ISS}$ at the observer; the uncertainty principle relates $\\tau_{\\rm ISS}$ to the bandwidth of the scintillations $\\Delta\\nu_{\\rm ISS} = 1/(2\\pi \\tau_{\\rm ISS})$. Over each element of the scintillation pattern $\\Delta\\nu_{\\rm ISS}\\times \\Delta t_{\\rm ISS}$, the propagation changes the electric field by a complex gain: a random amplitude and phase. At the source, the characteristic scale is $M S_{\\rm ISS}$; a displacement of the source by this length has an effect that is statistically equivalent to displacement of the observer by $S_{\\rm ISS}$. The above basic parameters define the arena of our ``interstellar telescope'': the high resolution of the scattering disk, acting as a lens, produces the statistical structure of the diffraction pattern in the observer plane, which is then modified by the structure of the source.  The Vela pulsar is particularly attractive for statistical studies of interstellar scattering because its scintillation bandwidth is relatively narrow at decimeter observing wavelengths, so that many samples can be accumulated quickly. The pulsar is strong, so the signal-to-noise ratio within one scintillation element is high. The pulsar is scattered enough that the diameter of the scattering disk can be measured with Earth-based interferometry; this measurement allows us to determine the location of the scattering screen along the line of sight, and also allows, in principle, studies of the two-dimensional structure of the source \\citep{Gwi01,Shi10}.  \\subsection{Pulsar Emission Structure via Interstellar Scattering}  A variety of authors have reported investigations of pulsar size using interstellar scattering. The observables address either motion of the emission centroid, or the size of the emission region. The present work falls into the second category.  Measurements of, or upper limits on, motion of the centroid of emission rely upon the reflex shift of the scintillation pattern in the observer plane, when the pulsar rotates \\citep{cor83,wol87,smi96,gup99}. % \\citep{wol87,smi96,gup99}. The observations measure such shifts from correlation of the scintillation spectrum at different phases of the pulse, with spectra at later or earlier times. Because motion of the source dominates changes in the scintillation pattern at the observer plane over these timescales, such a correlation, with information on the location of the screen, yields the shift of the source. These observations commonly find scales ranging from a few hundred to a few thousand km, or up to the diameter of the light cylinder.  The phrase ``stars twinkle, planets do not'' expresses the fact that a finite source emission region can decrease the modulation by scintillation. From the depth of modulation of scintillation, one can infer the size of the emission region \\citep{Coh66,Gwi98}. This technique was used before the advent of synthesis interferometry to measure source sizes from scintillation in the interplanetary medium, \\citep{Rea72,Hew74}. In ``strong'' scintillation, the modulation index $m=\\sqrt{{{\\langle I^2\\rangle}/{\\langle I\\rangle^2}} - 1}$ of a point source is 100\\% \\citep{coh74}; for most pulsars at decimeter wavelengths, scintillation is very strong. Thus, modulation of less than 100\\% would suggest the presence of source structure, on the scale $M S_{\\rm ISS}$. \\citet{Mac00} report an upper limit on the size of the Vela pulsar at $\\lambda=45\\ {\\rm cm}$ observing wavelength, from the modulation index, although they do not discuss effects of intrinsic variability of the pulsar or self-noise; these can also affect the modulation index.  We modify earlier approaches to measuring, or setting limits on, the size of the source from modulation, in that we fit a model to the distribution of intensity (or interferometric visibility). Source size affects the smallest intensities or visibilities most strongly. Intrinsic intensity variations and self-noise affect the largest intensities most strongly, as we discuss in Section\\ \\ref{sec:instrumental_effects}. All of these effects change the distribution function of intensity and visibility, and thus their moments.  The modulation index is a combination of first and second moments of intensity, and thus is most sensitive to the largest intensities. Thus, it is least sensitive to the shape of the distribution function where effects of source size are largest, and most sensitive where they are smallest. The full distributions of intensity or visibility provide more sensitive and complete information. They can distinguish among effects that alter these distributions in different ways.  \\subsection{Distributions of Electric Field, Intensity, and Visibility}  Scintillation affects the electric field of an astrophysical source at a particular frequency by a gain and a phase \\citep{Gwi11a}, here combined into the complex ``scintillation gain'' $\\tilde g$. For a point source in strong scattering, $\\tilde g$ is drawn from a circular Gaussian distribution in the complex plane, with zero mean. This behavior is a consequence of the large differences of path lengths, and the fact that many paths contribute to the signal received at the observer; the Central Limit theorem then implies, under rather general assumptions, that the scintillation gain resulting from that sum over paths is drawn from a Gaussian distribution. This result is independent of assumptions about the distribution of scattering material; for example, scattering in an extended or inhomogeneous medium can be treated quite generally via path integrals, and the same result holds \\citep{flatte}. In the time domain, the received signal is the emitted signal convolved with a kernel $g$ that describes the scattering medium; $g$ and $\\tilde g$ form a Fourier transform pair \\citep{Gwi11a}.  Because the scintillation gains are draws from a complex Gaussian distribution, their square modulus is drawn from an exponential distribution. An exponential distribution is thus the expected distribution of intensity for a scintillating point source, without effects of noise \\citep{sch68,Gwi98}. Similarly, the scintillation gains for two different antennas are drawn from correlated, complex Gaussian distributions. Their product, the interferometric visibility, is drawn from the distribution of the product of such quantities. This distribution is a zero-order modified Bessel function times an exponential \\citep{Gwi01}. In practice, to obtain these distributions the observer must average the intensity, or the interferometric visibility, over many samples of the random electric field of the source. This field is itself noiselike, and contributes to noise in the measurement via self-noise, as we discuss below.  If the source is extended, different points on the source will have different scintillation gains. These gains decorrelate as the separation between points on the source plane increases. A source consisting of two separated point sources provides a simple example. If the two parts are separated by much less than $M S_{\\rm ISS}$, then the gain factors are identical and the result is that for a point source above; if the separation is much greater than $M S_{\\rm ISS}$, then the observer records two superposed, independent scintillation patterns. If the sources are incoherent, then the observed intensity is the sum of the two; and the distribution of observed intensity is the convolution of two exponential distributions. The analogous results hold for interferometric visibility. For a small, but extended source, the distributions of gains and phases for the different parts of the source are correlated; however, they can be expressed as the convolution of the original distribution with distributions of the same form, but of smaller scales \\citep{Gwi01}. As discussed in more detail below, finite size tends to concentrate the distribution of intensity near the mean, and to soften the sharp cusp of the point-source distribution.  Noise  and intrinsic variations of flux density both broaden the observed distributions of intensity and interferometric visibility. Noise includes contributions both from backgrounds and from the noiselike source. Backgrounds are nearly independent of the flux density of the source, with small corrections for quantization effects \\citep{Gwi06,VelaNoise}. Source noise has standard deviation proportional to the flux density of the source; it is termed ``heteroscedastic,'' indicating that the variance is not constant \\citep{Osl11}. In combination, these contributions lead to variance of the noise given by a quadratic polynomial in phase with the signal, and the linear terms of that polynomial at quadrature \\citep{Gwi11b,VelaNoise}. Intrinsic variations of flux density can be divided into 3 regimes according to time scale \\citep{Gwi11b}. The time to accumulate one sample of the spectrum is the product of the sampling rate and the number of spectral channels, termed the ``accumulation time''. Variations shorter than the accumulation time introduce correlations but do not change the noise \\citep{Gwi11a}. Intermediate-term variations, longer than the accumulation time but shorter than the integration time, contribute to noise. Long-term variations, longer than the integration time, lead to a superposition of distributions with different mean flux density. Both noise and amplitude variations thus act in characteristic ways the can be distinguished from effects of size; moreover, both broaden the distribution rather than narrowing it, as does finite size.  \\subsection{Outline of Paper}  This paper focuses on our data and technique to estimate the size of the Vela pulsar's emission region from the distribution of interferometric visibility. This measurement requires data with stationary instrumental gain, well-characterized noise, and rapidly-sampled, gated correlation.  Our analysis involves models for the effects of scintillation, noise, and amplitude variations of the pulsar.  In Section\\ \\ref{sec:observations}, we describe our observations of the Vela pulsar, correlation and gating using the DRAO VLBI correlator, calibration, and fringing. We then present typical data and describe the formation of our data histograms.  Next, in Section\\ \\ref{sec:analysis}, we outline our analysis. This analysis involves calculation of model histograms and fits to data. We calculate the distribution of interferometric visibility in the complex plane, for a small, circular source, in Section\\ \\ref{sec:vis_dist}. We describe the distribution of noise in Section \\ref{sec:noise}, and demonstrate how a superposition of distributions can model the pulsar amplitude variations in Section \\ref{sec:amp_variations}. We discuss the combination of these effects and evaluation of the model in Section\\ \\ref{sec:numerical}. In particular, we demonstrate that the signature of finite source size is a W-shaped difference of the best-fitting finite-size model from the best-fitting zero-size model. We detail the numerical evaluation techniques in Section\\ \\ref{sec:model_calc}, and fitting techniques in Section\\ \\ref{sec:fitting}.  We present our results in Section\\ \\ref{sec:results_discussion}. We first discuss an example fit in detail and show that the data histogram displays the characteristic W-shaped signature of source size. We then give our results for all gates and spectral ranges. We discuss various systematic effects that can contribute to, or bias, the inferred source size. We briefly compare our results with previous results at $\\lambda=13$\\ cm, and compare with other observational studies of the Vela pulsar's emission region.  In Section\\ \\ref{sec:summary}, we summarize our results. ",
        "conclusions": "\\label{sec:results_discussion}  \\subsection{Finite- and Zero-Size Fits}\\label{sec:size_nosize}  We performed independent Levenberg-Marquardt fits for both finite and zero size sources. Comparison of the two provides a measure of the significance of our results. Both sets of fits were initialized with the best-fitting parameters for finite or zero size found from our grid search. Both allowed all the other parameters to vary: noise coefficients $\\{ b_0, b_1, b_2\\}$, amplitude scale $\\kappa_0$, and intrinsic amplitude variations on timescales longer than our integration time, $m_{\\rm s}$.  \\subsubsection{Sample Fit}  The lower panels of Figure\\ \\ref{fig:observed_distributions} show the residuals to a fit with zero source size, in histogram form. Structure remains in the residual histogram, apparent as variations near zero visibility. An independent fit, to the same data but with a finite source size, removes most of these variations. To illustrate this, we display the difference of the finite-size model from the zero-size model, as smooth curves. The difference curve for ${\\mathcal P}$ is the same as that shown in the lower panel of Figure\\ \\ref{fig:projected_model_distributions}. Like the residuals, this curve shows the W-shaped signature of finite source size.  The mean squared residuals in ${\\mathcal P}$, after fitting for finite source size, are approximately those expected for Poisson-noise dominated errors. The data contain 437 bins with more than 100 counts. The sum of the squared Poisson-weighted residuals, after fitting, is 434. Thus, the mean square errors are approximately as expected. The fit to a model for a source with zero size leads to a sum of the squared Poisson-weighted residuals of 621, significantly greater.  For ${\\mathcal Q}$, the sum of the squared residuals, after correction by the relative areas of ${\\mathcal P}$ and ${\\mathcal Q}$, and the factor of 3 from comparison of differences, is 705 for a finite-size model, and 785 for a zero-size model. The residuals errors in ${\\mathcal Q}$ appear noiselike in the plot, suggesting that the factor of 3 estimated from differences in other spectral and gate samples may be low for this sample of data. Nevertheless, the finite-size model improves the fit for ${\\mathcal Q}$.  The reduction in summed mean-square weighted residuals for ${\\mathcal P}$ and ${\\mathcal Q}$ together is 19\\%. This difference is highly statistically significant at more than the 40-$\\sigma$ level, for our sample size and number of parameters, according to the F-ratio test \\citep{Bev}. At this high level,  finite sampling limitations are unlikely to dominate errors in our estimates of emission size. However, this figure demonstrates that the signature of finite size appears in the data, as inspection of the figures suggests. The best-fitting size parameter is $(k M \\theta \\sigma)^2 = 0.0423$, or scaled size $(k M \\theta \\sigma) = 0.21$. This corresponds to a source size of approximately 180\\ km (standard deviation of the Gaussian distribution), or to approximately 420\\ km for the full width at half-maximum, as we discuss in Section\\ \\ref{sec:conversion_to_km} below.  Standard errors are small, because of the high significance of the fits as expressed by the F-ratio test. We present the best-fitting parameters for this IF, gate and channel range in Table\\ \\ref{tab:sample_parameters}. These values are typical for our fits. As we discuss in Sections\\ \\ref{sec:ceimf_size} and\\ \\ref{sec:systematic_effects} below, errors in the estimated size are probably dominated by systematic errors.   \\subsubsection{Fits to All Spectral and Gate Ranges}\\label{sec:ceimf_size}  We fit our model to spectral ranges of 1024 channels in each of the 5 gates, for both IFs. Our fits indicate that the pulsar has a rather large size at the start of the pulse, decreases in size over the first half of the pulse, and then increases in size again. Figure\\ \\ref{fig:pulseshape_size} shows the results of fits, giving the scaled size $(k M \\theta\\sigma)$ and mean amplitude $(\\kappa_0 + 2 \\kappa_1)$ as a function of pulse phase. Note that linear diameter is proportional to the square root of the fitted size parameter, $(k M \\theta\\sigma)^2$. All of the fits are independent; each point represents an independent sample of the pulsar's emission, fit completely independently using a priori parameters from independent grid searches (Section\\ \\ref{sec:gridsearch}). The two IFs are shown by different symbols (crosses for IF1 and circles for IF2), which represent distinct frequency ranges and thus completely different scintillation patterns. At some phases, points from different gates overlap; at these points, a higher frequency range in one gate coincides with a lower frequency range in the previous gate. Quite often the noise parameters for these different samples are quite different; nevertheless, size and amplitude track one another with pulse phase, despite the disparate origin of the samples. Some of the fits do sample identical samples of the diffraction pattern from interstellar scattering, as Figure\\ \\ref{fig:dynspec} suggests: however, these often lead to completely different results, indicating that the results arise from the pulsar, rather than from scattering. For example, data that includes the 3 panels shown in Figure 1 yields scaled sizes of $(k M \\theta\\sigma) = 0.08$, 0.09, and 0.40. (Of course, the grayscale in Figure\\ \\ref{fig:dynspec} emphasizes peaks of visibility, whereas the most critical information on source size is found near zero visibility, as Figure\\ \\ref{fig:projected_model_distributions} shows.) The dependence of fitted size on pulse phase, rather than on the frequency or details of the scintillation pattern, suggests that the effect arises from the pulsar rather than the scintillation pattern.  Our sample comprises 60 independent fits, over the two IFs, 5 gates, and 6 spectral ranges. A fit with finite size yielded significantly smaller residual for most of these: for the 43 fits with fitted size parameters of $(k M\\theta \\sigma)^2 > 0.017$, the reduction of the residuals was greater than 3\\%. This is significant at better than the 5-$\\sigma$ level according to the F-ratio test. Unsurprisingly, gate and spectral ranges with small fitted sizes, and regions with small amplitudes at the beginning and end of the pulse, show the least significant results. We do not display statistical errors from the fits on this figure; they are smaller than the plotting symbols.  The scatter of independent measurements at a given pulse phase provides a measure of the precision of our results. The residuals are dominated by variations among nearby bins, consistent with Poisson noise; whereas the models predict rather more slowly-varying differences across the histograms for ${\\mathcal P}$ and ${\\mathcal Q}$, as Figure\\ \\ref{fig:projected_model_distributions} would suggest. We made the bins narrow so as to follow the rapid variation of the histograms with ${\\rm Re}[V]$, as shown in Figure\\ \\ref{fig:observed_distributions}; we cannot broaden the bins without destroying this important information by averaging. We discuss averaging of the post-fit residuals over bins to visualize effects of the model in Section\\ \\ref{sec:pox} below.  \\subsubsection{Conversion to km}\\label{sec:conversion_to_km}  The angular width of the scattering disk $\\theta$ is important for the conversion of the fitted size parameter $(k M \\theta\\sigma)^2$ to $\\sigma$ in km. The angular width appears directly in this expression, and also affects the inferred magnification $M$ as we discuss below, so that the linear size of the emission region $\\sigma$ depends strongly upon $\\theta$. Unfortunately, our data do not provide a good determination of $\\rho$, as discussed in Section\\ \\ref{sec:rho} above, and so do not allow accurate determination of $\\theta$. Consequently, we adopt a simple conversion based on previous work, while noting that more accurate measurements of $\\theta$ and $M$ may revise the inferred conversion of $(k M \\theta\\sigma)$ to $\\sigma$, in kilometers.  \\citet{Gwi97} found angular broadening of $(3.3 \\times 2.2)\\  {\\rm mas}$ (full width at half-maximum intensity), with the major axis at position angle $92^{\\circ}$, in observations at wavelength $\\lambda=13\\ {\\rm cm}$. We adopt this value here. For simplicity in converting the observed $(k M \\theta\\sigma)$ to size $\\sigma$, we assume a circular scattering disk with full width at half maximum of $\\theta= 5.2\\ {\\rm mas}$; this is the mean square of the major and minor axes. For the scintillation bandwidth, we adopt the value of $\\Delta \\nu_s = 15$\\ kHz from Section\\ \\ref{sec:time_freq_averaging} below. With a distance to the pulsar of $D=287\\ {\\rm pc}$, as measured from the parallax of the pulsar \\citep{Dod03}, the combination of angular broadening and scintillation bandwidth yields a magnification of $M=D/R = 3.1$ \\citep{gwi93}. Consequently, the standard deviation of the Gaussian model for the source is $\\sigma_{\\theta} = (k M \\theta\\sigma)\\times 859\\ {\\rm km}$. In this expression, the subscript $\\theta$ serves to indicate the dependence of the inferred source diameter on the angular size of the scattering disk, and the magnification. The full width at half-maximum diameter of the fitted model is $\\sqrt{8 \\ln 2} \\sigma_{\\theta}$. The right-hand vertical axis on the lower panel in Figure\\ \\ref{fig:pulseshape_size} reflects these conversion factors. The size is as large as hundreds of km, expressed as standard deviation of a circular Gaussian distribution.  Changes in the assumed parameters for angular broadening will modify the inferred emission size. The angular scale of the scattering disk $\\theta$ sets the scale of the scintillation pattern at the source, and thus the resolution of the lens; and $\\theta$ sets the location of the scattering material along the line of sight, and so the magnification factor $M$. If the value assumed for $\\theta$ changes, the inferred size $\\sigma$ changes, even though the fitted parameter  $(k M \\theta\\sigma )$ remains unchanged. The net resulting effect is approximately proportional. Thus, halving the assumed angular broadening reduces the inferred size expressed in kilometers by a factor of 2, and so on. The previous measurement of $\\theta$ by \\citet{Gwi97} covered an incomplete arc in the $(u,v)$ plane, and depended on simpler approximations to the distribution of visibility for a scintillating source than the models used here; the measurement should be repeated. Refractive scintillation can change the angular broadening with time, but the variation is expected to be only about 85 microarcseconds for this line of sight \\citep{Rom86,Nar92}. If scattering material is distributed along the line of sight rather than in a thin screen, and some of it is located close to the source, it can increase the magnification $M$ without a change in the measured angular broadening. This would require scattering close to the pulsar, within the Vela supernova remnant, where density is expected to be low.  \\subsubsection{Binned Residuals}\\label{sec:pox}  We visualize the significance of our fits, by binning our residuals. This reduces effects of small-scale noise while allowing us to present the signatures of source size in the data. Philosophically, re-binning the residuals is similar to subtracting a model from the data. For example, we subtract a model for source position, baseline length, and Earth orientation from the phases in the correlator, so that the ``sky'' fringe rate can be reduced enough for time integration to recover the fringe with high signal-to-noise ratio. Fringing reveals the inaccuracies of the subtracted model. Here, analogously, we remove a model for the distribution of visibility from the histograms of data, reducing variations between neighboring bins from noise. We then average the residuals of nearby bins so that the effects of source structure are more easily apparent. This averaging reveals the degree to which our model fits match the observations.  Figure\\ \\ref{fig:poxres} shows the result of binning residuals, again for the data in IF1, Gate 1, channels 4096 to 5120. After fitting zero-size and finite-size models, the residuals were binned and averaged in groups of 20, for resulting widths of $20\\, w = 0.0040$ in ${\\rm Re}[V]$. Averaging (rather than summing) keeps the normalization the same. The dotted histogram shows the residuals from the fit with zero size for the source. We also display the differences between the model for finite source size and for zero size as solid lines. The curves are the same as those shown in Figures\\  and\\ \\ref{fig:observed_distributions} and \\ref{fig:projected_model_distributions}. A zero-size source would show a flat, dotted histogram. A model that perfectly corrected the deficiencies of the zero-size model would track the dotted histogram perfectly, allowing for the finite widths of the bins. Clearly, adding a size parameter explains most of the slowly-varying residuals. Quantitatively, after binning the mean square residual is reduced by a factor of 13, by including a size parameter. This is much larger than the reduction observed before binning. As Figure\\ \\ref{fig:observed_distributions} suggests, Poisson noise dominates the residuals without binning, and is greatly reduced by binning.  The model fit is imperfect. Some systematic variations remain, more so in some ranges of pulse phase than others. Figure \\ref{fig:triplepox} shows residuals for the zero-size model, and differences between finite-size and zero-size models, for representative spectral ranges in three gates. Fits to the same spectral range and same gate, but different IF, resemble one another strongly, as expected because they lead to nearly the same fitted size. Noise and amplitude parameters vary between IFs, and even more between gates and channel ranges, so that the binned residuals are not identical.  The upper pair of panels of Figure \\ref{fig:triplepox} show IF2, gate 1, channels 4096 to 5120: these data are equivalent to those in Figure\\ \\ref{fig:poxres} in pulse phase and spectral range, but are for the other IF. The data are thus completely independent. Introduction of a size parameter reduces the mean square residual of $\\mathcal P$ by 92\\%, and $\\mathcal Q$ by 12\\%. The best-fitting size parameter is $\\kappa_1/\\kappa_0 = (k M \\theta \\sigma)^2 = 0.043$, or scaled size $(k M \\theta \\sigma) = 0.21$. The IFs agree closely in residuals and in fitted size. This range lies on the rising part of the pulse profile, before the peak.  The middle pair of panels show IF1, gate 2, channels 2048 to 3072. This range lies near the peak of the pulse. The best-fitting size parameter is $\\kappa_1/\\kappa_0 = (k M \\theta \\sigma)^2 = 0.006$, among the smallest results we obtain. Before binning, introduction of the size parameter reduced the Poisson-weighted residuals by less than 1\\%; this reduction is significant at the 2-$\\sigma$ level according to the F-test. After binning, the reduction in mean square residual for ${\\mathcal P}$ is 32\\%; however,  zero-size and finite-size models both fit $\\mathcal Q$ almost equally well. Some systematic effects appear, particularly the sharp downward spike in $\\mathcal P$ near ${\\rm Re}[V]=0$.  The lower pair of panels show IF2, gate 3, channels 4096 to 5120. This range lies on the trailing side of the pulse, near where it briefly flattens into a plateau. At shorter observing wavelength, this plateau becomes a second component that arises close to this pulse phase. Introduction of a size parameter decreases the residuals,  by 10\\% in $\\mathcal P$ and by 4\\% in $\\mathcal Q$ before binning, and by 54\\% and 24\\% after binning. Again the influence of $\\mathcal Q$ on the overall fit is small, because zero-size and finite-size models both fit $\\mathcal Q$ well. The best-fitting size parameter is $(k M \\theta \\sigma)^2 = 0.039$, or $(k M \\theta \\sigma) = 0.20$. Although the fit to $\\mathcal P$ is good, some systematic variation appears as a shifting of the histogram relative to the best-fitting model; our circular-Gaussian model does not accommodate such a shift. We are investigating a variety of simple models for the systematic variations, as we discuss briefly in Section\\ \\ref{sec:alternative} below.  \\subsection{Systematic Effects}\\label{sec:systematic_effects}  The effects of size that we observe in our data include its W-shaped signature in the distribution of visibility ${\\mathcal P}$, its variation with pulsar phase, and its constancy for different spectral ranges and samples of the scintillation pattern at the same pulse phase. The W-shaped signature of finite size on the distribution of visibility is quite characteristic, as Figures\\ \\ref{fig:observed_distributions},\\ \\ref{fig:projected_model_distributions}, and \\ref{fig:poxres} suggest, and appears for both IFs and many pulse gates and spectral ranges. The agreement of fitted size is good between IFs, for overlapping spectral ranges at different frequency but the same pulse phase, and between nearby but different spectral ranges. Effects that can match all of these, or even only the latter two, are nearly all associated with geometrical effects at the pulsar.  \\subsubsection{Errors in Model Parameters}  Further significant changes in the magnification factor from changes in pulsar distance seem unlikely. The distance from parallax is far more reliable than that from dispersion measure. In principle a model that included significant scattering in the immediate neighborhood of the pulsar, as well as scattering in the surrounding Vela supernova remnant, could increase the effective magnification by bringing the lens closer to the pulsar. Temporal broadening at low frequencies suggests that scattering material is concentrated into a thin screen \\citep{Joh12c}. Previous measurements of the angular broadening $\\theta$ should be revisited, as described in Section\\ \\ref{sec:conversion_to_km} above. Further studies of the angular broadening, using long baselines with different orientations, can help to improve on the earlier results. This may change the inferred linear size $\\sigma$.  \\subsubsection{Instrumental Effects}\\label{sec:instrumental_effects}  Instrumental effects are unlikely to reproduce the signature of source size on the distribution of visibility, as shown in Figure\\ \\ref{fig:poxres}, because they tend to vary with pulse gate and frequency within the passband. For example, for both IFs, the noise parameter $b_0$ shows a similar pattern for each pulse gate, reflecting variations in the gain and noise in the passband, with an overall offset for each gate that reflects changes in quantization noise \\citep[see][Section 5.1]{VelaNoise}. Effects of quantization, in particular, can be expressed in the spectral domain as a gain and a change in noise, and so are absorbed into the amplitude parameter and $b_0$. Effects of variation of noise parameters would be expected to appear as differences between overlapping gates.  The direction of linear polarization varies smoothly across the pulse. We observe only left-circular polarization, the fraction of which varies little. Scintillation in the interstellar medium is nearly polarization-independent \\citep{Min96,Spa01}. Defects in separation of polarizations at the antennas could produce artifacts that vary with pulse phase, but these can be represented accurately as pulse-phase-dependent gains. They would not produce the distortion of the distribution of visibilities that we observe as the signature of source size.  Saturation effects can be important at large antennas such as Tidbinbilla. For our observations, these might affect particularly strong pulses, such as the giant pulses reported by \\citet{Joh01}. However, such pulses carry an insignificant fraction of the total flux density of the pulsar. Saturation effects can also reduce visibility for the very strongest scintillations, at least for baselines involving two large antennas \\citep[][Section 4.5]{Gwi00}. These effects are expected to distort the distribution of visibility at the largest values, but are not likely to reproduce the W-shaped signature of finite size we observe.  \\subsubsection{Calibration and Fringing Effects}  We perform relatively little calibration on the data.  We do not change amplitude, particularly within a scan or with frequency. We do change phase, in fringing. Mis-fringing is possible for observations of a scintillating pulsar in the speckle limit, because the phase and amplitude vary on the frequency and time scales of scintillation \\citep{des92}, so that the model used for fringing can be challenging to fit. Mis-fringing would tend to increase phase variations, while keeping amplitude the same. Such effects are reduced for a short baseline, such as the Mopra-Tidbinbilla baseline, because the phase variation is less. If the phase errors were very large, it could potentially alter the distribution of visibilities in our projections, by altering the relative sizes of real and imaginary part. We reduced errors from mis-fringing by fitting our fringe model over approximately 2 decorrelation times and 800 scintillation bandwidths. The wide bandwidth reduces effects of scintillation-based phase variations. The short time was chosen to match any possible ionosphere- or atmosphere-induced phase variations.  Effects of errors in the fringe model on size would be expected to reproduce among gates, since the fringe model from Gate 2 was applied to all gates (with the exception of a single phase across the band, fitted for each gate); however, size varies with pulse phase, rather than by channel range within gates, and is near zero for some spectral ranges near the pulse peak. For channel ranges with sufficient signal-to-noise ratio, such as the later channels of Gate 1 and most channels of Gate 3, we found that fringing to the data within that gate, rather than using the model from Gate 2, produced identical results. Similarly, applying the model from other strong gates to Gate 2 yielded nearly the same results.  \\subsubsection{Time and Frequency Averaging}\\label{sec:time_freq_averaging}  Effects of the finite sampling of the diffraction pattern are small, and are expected to be constant across the pulse. Sampling over finite intervals of time and frequency blurs together correlated samples of the diffraction pattern, and so is indistinguishable from effects of source size \\citep{Joh12a}. For given scintillation timescale $t_{\\rm ISS}$ and scintillation bandwidth $\\Delta\\nu$, effects of sampling can be calculated analytically \\citep[][Section 2.2]{Gwi00}.  We estimate the scintillation timescale and bandwidth from our observations on the Mopra-Tidbinbilla baseline, using data near the pulse peak, IF2, Gate1, channels 2048 to 3072. We find correlation functions for the observed real part of interferometric visibility, and fit expected functional forms to them \\citep[see][Equations 5, 7]{Gwi00}. Figure\\ \\ref{fig:scintillationscales} displays the data  and results. We have normalized the correlation functions using the amplitude and offset of the best-fitting functional forms. The data point at zero lag is omitted, because it is elevated by effects of noise and variations of the amplitude with time. Both correlation functions extend to large lag, off the right of the displayed panels; these are included in the fits, but not displayed here.  A fit of the expected Gaussian function to the temporal correlation function yields a scintillation timescale of $t_{\\rm ISS}= 8.99\\pm 0.17$\\ s. Our sampling time of 2\\ sec is then expected to decrease the mean square flux density by a factor of $f_{t} = 0.994$ \\citep[][Equation 8]{Gwi00}, while leaving the average flux density unchanged. The variation of $f_{t}$ from the standard error of the fit is less than 0.001.  A fit of the expected Lorentzian function to the frequency correlation function yields a scintillation bandwidth of $\\Delta\\nu = 15.2\\pm 0.5$\\ kHz. Our channel bandwidth of 1.95\\ kHz is then expected to decrease the mean square flux density by a factor of $f_{f} = 0.998$ \\citep[][Equation 6]{Gwi00}, leaving the mean flux density unchanged. The observed correlation function shows systematic departures from the best-fitting Lorentzian function, with a sharper peak and elevated correlation at intermediate lag, with a peak near a lag of 70\\ kHz. A Kolmogorov rather than a square-law structure function, and a departure of the scattering disk from isotropy, can produce such effects\\ \\citep{Gwi01,Joh12a}. Temporal variation of the emission on short timescales will also introduce correlations in frequency\\ \\citep{Gwi11a}. Models including such effects will fit better. However, as the figure shows, the correlation falls to half its peak value at a lag of approximately 14\\ kHz, providing a clear characteristic scale. The variation of $f_{f}$ from the standard error of $t_{\\rm ISS}$ is less than 0.001. Even with $\\Delta\\nu=12$\\ kHz we still find $f_{f}=0.998$.  Together, the effects of finite sampling in time and in frequency will increase the fitted size. Under the assumption that the source is pointlike, so that $(kM\\theta\\sigma)=0$, and we misinterpret the effects of averaging in time and frequency as source size, our sampling parameters lead to a fictitious size parameter of $(k M \\theta \\sigma) = 0.06$, where we have related effects of finite sampling to size using the expressions for modulation index from averaging \\citet[][Equation 9]{Gwi00} and for source size \\citet[][Equations 34, 35]{Gwi01}. This value is indeed nearly the minimum that we observe, near the pulse peak, as Figure\\ \\ref{fig:pulseshape_size} shows.  The effects of finite sampling in time and frequency are constant with pulse phase, for a pointlike source. This is clearly inconsistent with our results discussed in Section\\ \\ref{sec:results_discussion} above. Variation of amplitude with time can affect estimates of the scales of scintillation: variability on long timescales can affect the correlation with time, and on short timescales can affect the correlation in frequency \\citep{Gwi11a,Gwi11b}. Erroneous estimates will not produce variations of source size with pulse phase. Direct effects of source variability on estimates of source size are discussed elsewhere in this paper (Sections\\ \\ref{sec:amp_variations},\\ \\ref{sec:noiseandamplitudevariations}). %Variation of amplitude with time between samples can reduce the correlation of samples in time. %However, unless the variability is correlated over times comparable to the scintillation timescale, %the effect will appear completely at the zero lag, which we omit from our analysis. %Variability is nearly uncorrelated between successive pulses, and falls rapidly with pulse separation, whereas we average over groups of 22 or 23 pulses. %We expect the effect to be of order %Variation of amplitude with time over times shorter than the accumulation time introduces correlations between spectral channels, %and so acts to convolve the true spectral correlation function with variability kernel \\citep{Gwi11a}. %This broadens the correlation function in frequency, and so increases the inferred value of $f_{f}$, without an accompanying effect on %the source size we infer. %Thus, source variability on short timescales leads to an overestimate of the effects of sampling on size.  \\subsection{Alternative Models}\\label{sec:alternative}  Our fit assumed a circular-Gaussian distribution of emission at the source. More precisely, since the scattering disk is presumably elongated, we have fit the data with a model for a source with the same axial ratio as the scattering disk. This description simplifies the mathematics. Some of the residuals in Figure\\ \\ref{fig:triplepox} show systematic residuals that suggest that a more complicated model might provide a better fit. A source elongated along a single direction provides a model arguably as simple as the one we use; the size inferred for such a fit is approximately $\\sqrt{2}$ greater than the diameter for a circular Gaussian. A source of arbitrary axial ratio requires a more complicated model; we will discuss such models in conjunction with data on longer baselines. Of course, an infinite set of models can be fit to any finite data set; we choose the circular-Gaussian case as providing a simple parameterization that agrees well with our data.  Another simple model is a ``core-halo'' model, including a pointlike core superposed with an extended halo. In the simplest case, this model would include a pointlike core and a halo so extended that it did not scintillate at all. Such a model would produce the distribution of visibility for a scintillating point source, as shown for example in the left panel of Figure\\ \\ref{fig:size_nosize_nonoise_distribution}, offset to the right, by convolution with the delta-function distribution of visibility expected for a non-scintillating source. To match our observations, the two components of this distribution would have to change their relative magnitudes over the course of the pulse, and to both disappear when the pulsar is ``off\": we see neither in the empty Gate 6. Moreover, the size of the ``halo'' would have to be large: much greater than 800\\ km. We regard this model as more complicated, and less natural, than a circular-Gaussian model. Such a model appears to fit less well than a circular-Gaussian model at either side of the pulse maximum, where size and amplitude are both relatively large, although it fits significantly better than a zero-size model at some ranges of pulse phase.  \\subsection{Comparison with $\\lambda = 13$ cm results}  We reported a size of the Vela pulsar's emission region of approximately 300 km, at 13\\ cm observing wavelength, in previous work \\citep{Gwi97,Gwi00}. The three gates used there correspond roughly to ranges of pulse phase of $-1.2$ to 0.2\\  msec, 0.2 to 1.2\\ msec, and 1.2 to 3.7 msec, in Figure\\ \\ref{fig:pulseshape_size}. We reported size parameters of $(k M \\theta \\sigma)^2$ of $0.091\\pm 0.009$, $0.070\\pm 0.009$, and $ 0.020\\pm 0.020$ in three gates across the pulse, at $\\lambda=13\\ {\\rm cm}$ observing wavelength. (Note: The exponent ``2'' was omitted in the legend for these quantities in the first line of the body of Table 4 of \\citet{Gwi00}). For comparison with the lower panel of Figure\\ \\ref{fig:pulseshape_size}, the square roots of these measured quantities should be multiplied by the ratio of wavelengths of the observations, 18/13. The resulting sizes are comparable to those inferred here, except for the third gate, for which the size at 13\\ cm wavelength is significantly smaller. The 13-cm gate extends well past the end of the pulse at 18 cm; the pulse is wider at the shorter wavelength, and flux density in the gate is significantly higher, because of the emergence of a second peak.  This profile difference might be responsible for the emission size discrepancies.  The observations and analysis at 13-cm wavelength differed in detail from those presented here. One difference was that we assumed a circular source, rather than one with the same aspect ratio as the scattering disk; this would tend to increase the size parameter $(k M \\theta \\sigma)^2$ estimated at 13\\ cm wavelength. We also used an exponential model; this is a reasonable approximation for the shorter baseline at shorter wavelength. The revised distance to the pulsar from parallax measurement, of  $D=287\\ {\\rm pc}$ \\citep{Dod03} rather than $500\\ {\\rm pc}$, changes the magnification factor $M$ from $1.5$ to $3.1$ and so decreases the estimated size of the pulsar. The integration time was longer and the spectral-channel bandwidth was higher, as allowed by the increased decorrelation scales of the scintillation. However, the channel bandwidth was great enough that the effects of averaging in frequency were non-negligible \\citep[see][Section 2.2]{Gwi00}. There were also fewer scintillation elements sampled. The model was fit to amplitude rather than to interferometric visibility; consequently, the effect of size on the distribution function is less simple. The model for self-noise omitted the quadratic term, although this omission was justified because the effects of finite source size appear at small amplitude, where the quadratic term is small.  A 2-dimensional distribution of noise as given by $b_0$ and $b_1$ is critical for an accurate fit to the distribution of amplitude, however.   \\subsection{Comparison with Pulsar Geometry}\\label{sec:geometry}  The pulse profiles of most pulsars can be divided into ``core'' and ``cone'' components, after the inferred shape of the beam \\citep{ran83,ran90}. Vela is usually classified as a ``core single\" pulsar \\citep{ran93}, as might be expected from its young age and short period; its narrow pulse, with width nearly constant with observing wavelength; and its strong polarization at meter wavelengths. A second component emerges at observing wavelength $\\lambda \\sim 10$\\ cm, although it is not clear that this represents a subsidiary peak \\citep{Ker00}.  % ; refs from kern et al Vela fits the scaling relation of pulse width with period for conal pulsars well \\citep{ran93}. This scaling matches the opening angle for a dipole field near the star's surface, suggesting that core-single emission arises from near the star's surface. The radio beam is not well aligned with the magnetic axes suggested by high-energy emission, as is not uncommon for single-core components, raising the possibility of other geometry for this type of emission.  Unlike most core-single pulsars, Vela shows an organized pattern of polarization; indeed, it was the archetype of the rotating-vector model, which provides the fundamental pattern for conal emission \\citep{rad69,kom70}. The angle between the magnetic axis and the line of sight is $6^{\\circ}$ in this model \\citep{Joh01}; if one assumes that emission arises where the last open magnetic field lines are tangent to the line of sight, and that the rotation axis lies in the plane of the sky, the altitude of emission is 200\\ km. The inferred altitude is significantly greater if the rotation axis is not in the plane of the sky.  \\citet{kri83} analyzed the polarization properties of the Vela pulsar, in bins in amplitude of individual pulses. They used the mapping of polarization onto dipole field geometry, and the difference in arrival time, to infer the location of four distinct pulse components. They found that earlier emission components arise from higher above the neutron-star surface than later components, with a spread in altitudes of 500\\ km.  In contrast, our measurements are sensitive to the lateral extent of the emission region, rather than its altitude. These two are related by the height of the emission region and, plausibly, the geometry of the magnetic field. Detailed discussion is beyond the scope of this paper, but we note that the inferred lateral sizes tend to suggest emission altitudes at least as great as those inferred from field geometry \\citep[][and references therein]{Dyk04,Gan05,Joh12b}.  The observed size might result from refraction or scattering in, or near, the pulsar's emission region. The last closed field lines provide zones where plasma waves might grow, and then then propagate along the field until converted to electromagnetic waves \\citep{Bar86}. Such an emission process might preserve much of the geometry and temporal variations of the emission, while translating it to a higher altitude \\citep{hir01}. Interestingly, polarization appears to be perpendicular to the curvature of the magnetic field lines for the Vela pulsar, which suggests reprocessing after emission \\citep{Lai01}. Individual samples of the electric field of the pulsar show log-normal statistics, suggestive of multiple, random contributions to amplification \\citep{cai03}. Radiation emitted at low altitudes may also be refracted or scattered at higher altitude \\citep{lyu00}. Reprocessing might include interaction with X-ray and gamma-ray emission \\citep{har08,pet09}. Thus, a variety of evidence suggests that the site of the initial emission may be distinct from the location at which the radiation is released to propagate, with scattering in the interstellar medium, to the Earth.     We have described analysis of the distribution of visibility of the Vela pulsar at $\\lambda=18$\\ cm wavelength. We observed the pulsar in two IFs of 16\\ MHz bandwidth each.  We observed the pulsar in 5 gates across the pulse, each 1\\ msec wide. The signal was not dedispersed, so higher frequencies sampled later pulse phase than lower frequencies. We fringed the data and from it formed distributions of visibility projected into bins along the real axis ${\\mathcal P}\\left( {\\rm Re}[V]\\right)$ and the second moment of imaginary part in each bin ${\\mathcal Q}\\left( {\\rm Re}[V]\\right)$.  We calculated a theoretical model for the distribution of visibility.  This model used the distribution of interferometric visibility for a scintillating source \\citep{Gwi01}. We calculated models for a source of zero size, and for a source of small but nonzero size. We assumed a Gaussian distribution of flux density with the same axial ratio as the scattering disk.  We then added the effects of noise and self-noise using convolution with a non-stationary kernel. We modeled noise with a second-degree polynomial in phase with source visibility, and the linear terms of that polynomial at quadrature, as predicted by theoretical models for noise and exhibited by our data \\citep{VelaNoise}. We also included a parameter for effects of intrinsic amplitude variability of the source, on timescales longer than the 2-sec integration time. We find that effects of size appear as a W-shaped difference of the best finite-size model from the best zero-size model, as a consequence of the shape of the underlying distributions of visibility.  We fit our models to the observed distributions using the Levenberg-Marquardt method. We weighted residuals by Poisson statistics for ${\\mathcal P}$ and by statistics appropriate for mean square elements of a Gaussian distribution for ${\\mathcal Q}$, with constant weight below a cutoff of $N\\leq 100$ samples in a bin. We searched parameter space using a grid search with the exponential $\\rho=1$ form for the visibility distribution, and found that the variations of the summed, squared weighted residuals was quite simple. We used the minimum parameters from this search as initial parameters for Levenberg-Marquardt iterative fits.  Residuals to our fits for zero size show the characteristic W-shaped signature of finite size.  Independent fits for finite size describe the data well, matching the data and presenting the expected form when differenced with the zero-size model. We find that the resulting inferred, intrinsic size of the pulsar emission region first decreases across the pulse, then increases.  The maximum size is approximately 800\\ km (FWHM of our Gaussian distribution), and the minimum is near zero. The reduction in residuals ranges up to 10\\%. The statistical significance of including the size parameter is high, 40-$\\sigma$ or more in some cases. Residuals are dominated by Poisson noise, from the finite number of samples in each bin.  To help visualize effects of the fits, we binned the residuals after fitting. Binning reduces Poisson noise, and removes the relatively rapid variation of the distribution, leaving the slower variations from source size. The results show that the distribution of visibility matches the W-shaped signature of finite source size well, as we observe in a test case even before binning. After removing most Poisson noise by binning, introduction of the fitted size parameter reduces residuals by as little at 30\\% and as much as 92\\%. Some gates show evidence for residual systematic differences of the distribution from the model. We consider various systematic effects that could modify our results, and compare results with previous observations at $\\lambda=13$\\ cm observing wavelength. We briefly compare our results with previous work on geometry of pulsar emission."
    },
    "1208/1208.1328_arXiv.txt": {
        "abstract": "The excited state of the $^{12}$C nucleus known as the ``Hoyle state'' constitutes one of the most interesting, difficult and timely challenges in nuclear physics, as it plays a key role in the production of carbon via fusion of three alpha particles in red giant stars. In this letter, we present \\textit{ab initio} lattice calculations which unravel the structure of the Hoyle state, along with evidence for a low-lying \\mbox{spin-2} rotational excitation. For the $^{12}$C ground state and the first excited \\mbox{spin-2} state, we find a compact triangular configuration of alpha clusters. For the Hoyle state and the second excited \\mbox{spin-2} state, we find a ``bent-arm'' or obtuse triangular configuration of alpha clusters. We also calculate the electromagnetic transition rates between the low-lying states of $^{12}$C. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.4950_arXiv.txt": {
        "abstract": "We investigate the occurrence of crystalline silicates in oxygen-rich evolved stars across a range of metallicities and mass-loss rates. It has been suggested that the crystalline silicate feature strength increases with increasing mass-loss rate, implying a correlation between lattice structure and wind density. To test this, we analyse {\\em Spitzer} IRS and {\\em Infrared Space Observatory} SWS spectra of 217 oxygen-rich asymptotic giant branch and 98 red supergiants in the Milky Way, the Large and Small Magellanic Clouds and Galactic globular clusters. These encompass a range of spectral morphologies from the spectrally-rich which exhibit a wealth of crystalline and amorphous silicate features to `naked' (dust-free) stars. We combine spectroscopic and photometric observations with the {\\sc grams} grid of radiative transfer models to derive (dust) mass-loss rates and temperature. We then measure the strength of the crystalline silicate bands at 23, 28 and 33 $\\mu$m. We detect crystalline silicates in stars with dust mass-loss rates which span over 3 dex, down to rates of $\\sim$10$^{-9}$ M$_\\odot$ yr$^{-1}$. Detections of crystalline silicates are more prevalent in higher mass-loss rate objects, though the highest mass-loss rate objects do not show the 23-\\mum feature, possibly due to the low temperature of the forsterite grains or it may indicate that the 23-\\mum band is going into absorption due to high column density. Furthermore, we detect a change in the crystalline silicate mineralogy with metallicity, with enstatite seen increasingly at low metallicity. ",
        "introduction": "Cool evolved stars in the later stages of active nuclear burning expel a large fraction of their mass through a slow, dense stellar wind at rates of $10^{-7}$ to $10^{-4}$ \\Msuny \\citep{Habingbook}. In low- to intermediate-mass stars (0.8--8~M$_\\odot$) this occurs on the asymptotic giant branch (AGB), while higher-mass stars (8--25 \\Msun) undergo this intense mass loss as red supergiants (RSG). In AGB stars it is believed that radial pulsations levitate gaseous material above the surface of the star \\citep{Habingbook}, resulting in a cooling outflow from which molecular species form and (at temperatures below $\\sim$1400 K) dust grains condense \\citep{Hoefner1998}. A dust-driven wind develops, propelled by radiation pressure from the stellar photons  \\citep{Gehrz1971, Bowen1991, Norris2012}. For RSGs, a slightly different mechanism may drive the dense outflow \\citep{JosselinPlez2007}.  Chemistry in AGB atmospheres is affected by mixing of the stellar atmosphere and interior, known as dredge up. This process brings nuclear-fusion products, particularly carbon, to the stellar surface, enhancing their abundances. Chemistry in the cooling wind is dominated by the relative abundances of carbon and oxygen, as these combine at high temperatures to form carbon monoxide. The remaining free carbon or oxygen then determines whether the star produces C-rich dust (such as amorphous carbon and SiC) or O-rich dust (such as silicates and metal oxides). As we are interested in silicate dust production, we focus this paper only on O-rich stars.  Infrared (IR) spectroscopic observations of O-rich evolved stars show a plethora of features due to molecules and dust \\citep{Omont1993, Waters1996, Sylvester1999, Molster2002a}. Molecules, mainly CO, OH, H$_{2}$O, and SiO, dominate the spectra at $\\lambda \\lesssim 10 \\mu$m. At longer wavelengths, dust features become more important. Silicate dust dominates the IR spectra of the majority of dust-producing O-rich AGB stars, and can be found in either amorphous or crystalline form. Amorphous silicates are most abundant, characterised by their broad, smooth features at 10 and 18 $\\mu$m. Crystalline silicates exhibit sharp resonance features at 10 $\\mu$m and longer wavelengths. The positions and shapes of these are very sensitive to compositional changes, lattice structure, and grain size and morphology, providing a mechanism to identify the chemical composition and mineralogy of the dust grains \\citep{Molster2002b, Chihara2002, Koike2003, Min2003, Molster2005}. Comparisons with laboratory data have  primarily identified these grains as Mg-rich olivines and pyroxenes (forsterite: Mg$_{2}$SiO$_{4}$ and enstatite: MgSiO$_{3}$) with little or no iron content \\citep{Molster1999, deVries2011}.  The formation of crystalline silicates in the wind of evolved stars is not well understood: crystalline structures are the energetically more favourable atomic arrangement for silicates yet, in most cases, crystalline silicate dust is only a minor component of the circumstellar dust shell. The transition from amorphous to crystalline silicate grains is a thermal process, requiring temperatures in the region of 1040 K \\citep{Hallenbeck1998, Fabian2000, Speck2011}. However, there is little consensus on how silicates in stellar outflows gain sufficient energy for this transition to occur. In order to determine the physical conditions under which crystalline silicates form we require a better understanding of the conditions surrounding the star (e.g. wind densities). The formation process of crystalline grains can be constrained by correlating the crystalline fraction with the dust or gas column density. A correlation with dust density suggests that annealing of amorphous silicate grains heated by radiation is probably the dominant means by which crystals are manufactured \\citep{Sogawa1999}, while a correlation with gas density would suggest that direct condensation of crystalline silicates in the wind will dominate \\citep[e.g.][]{Tielens1998, GailSed1999}.  In {\\em Infrared Space Observatory} ({\\em ISO}) spectra of Milky Way (MW) giant stars, the spectral features due to the crystalline silicates forsterite (Mg$_{2}$SiO$_{4}$) and enstatite (MgSiO$_{3}$) typically only appear around evolved stars if their mass-loss rate is higher than a threshold value of $\\sim$10$^{-5}$ M$_\\odot$ yr$^{-1}$ \\citep{Cami1998, Sylvester1999}. For example, the high-density winds of heavily enshrouded OH/IR stars have measured crystalline fractions of up to $\\sim$20 per cent of the silicate mass \\citep{Kemper2001, deVries2010}. Conversely, lower mass-loss rate AGB stars, such as Miras, normally lack crystalline silicate dust features in their spectra. This suggests a correlation between lattice order and wind density (e.g. \\citealt{Cami1998, Speck2008}). Alternatively, contrast effects between amorphous and crystalline silicate grains at different temperatures can mask the characteristic spectral features of the crystalline material (up to the 40 per cent mass fraction level) in infrared spectra of low mass-loss rate AGB stars  \\citep{Kemper2001}. If the crystalline and amorphous silicates are in thermal contact with each other, contrast improves, and the detection of smaller amounts ($<$40 per cent) of crystalline silicates becomes possible. This could explain observations of crystalline silicates in lower mass-loss rate objects, such as the recent observations of low-metallicity, low mass-loss rate evolved stars in Galactic globular clusters that show crystalline silicates \\citep{Sloan2010,McDonald2011a,Lebzelter2006}.  \\begin{table} \\begin{minipage}{84mm} \\caption{LMC and SMC Parameters} \\label{tab:galProperties} \\begin{tabular}{@{}lcccc@{}} \\hline \\hline Parameter    &   LMC    &      Ref.     &      SMC    &    Ref. \\\\ \\hline Distance,    $d \\: (\\rm {kpc})     $   &  $ 51 \\pm 2    $  &  $   1,2    $ &    $ 61 \\pm 2    $    &   $  1,2    $ \\\\ Metallicity, $Z \\: (Z_{\\odot})      $  &  $ 0.5 \\pm 0.17 $ &  $   3,4    $ &    $ 0.2 \\pm 0.06$    &   $  3,4    $ \\\\ Inclination angle  $( ^{\\circ} ) $     &  $ 34.7 \\pm 6.2 $ &  $   5      $ &    $ 68 \\pm 2    $    &   $    6    $ \\\\ Gas-to-dust ratio  $ ({\\it \\Psi}) $    &  $ 200        $   &    $\\ldots$   &    $ 500       $      &   $\\ldots   $ \\\\ E(B-V) (mag)                           &  $ 0.13         $ &  $   7      $ &    $ 0.04        $    &   $   8     $ \\\\ \\hline \\multicolumn{5}{p{0.95\\textwidth}}{ {\\sc References:} (1) \\citet{Cioni2000}; (2) \\citet{Szewczyk2009}; (3) \\citet{Luck1998}; (4) \\citet{Meixner2010}; (5) \\citet{vanDerMarel2001}; (6)  \\citet{Groenewegen2000}; (7) \\citet{Massey1995}; (8) \\citet{Harris2004}.} \\end{tabular} \\end{minipage} \\end{table}   Observations of evolved stars in the metal-poor environments of the Magellanic Clouds with the {\\em Spitzer Space Telescope} provide an ideal opportunity to explore the occurrence of crystallinity and investigate how the O-rich dust condensation depends on the physical and chemical conditions of the envelope. Within the Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC), we have a set of stars with similar characteristics: for instance, they are all at approximately the same distance from the Sun and are assumed to have a single metallicity within each galaxy. This allows us to measure luminosities and dust mass-loss rates from the observed spectral energy distributions. The parameters adopted in our calculations for the Magellanic Cloud stars are listed in Table \\ref{tab:galProperties}. The dust-to-gas mass ratios of stars in the Magellanic Clouds are assumed to be lower than those of stars in the Solar neighbourhood and have a linear dependence with metallicity \\citep{vanLoon2000,Marshall2004,vanLoon2006}, breaking the degeneracy between the dust column density and the gas density in the outflows of AGB stars. By carefully studying the dependence of crystallinity on the dust and gas mass-loss rate, we address the influences of dust density and gas density on the formation of crystalline grains.  This paper is organised as follows: Section \\ref{sec:sample} describes the samples. In Section~\\ref{sec:analysis} we fit the spectral energy distribution (SED) of each source and measure the strength of crystalline silicate emission bands. Average feature profiles are shown in Section~\\ref{sec:results}, and the results are presented. Finally in Section~\\ref{sec:discussion} we discuss the implications for the onset of crystallinity and the change in mineralogy with metallicity. The conclusions are summarised in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have analysed \\emph{Spitzer} and \\emph{ISO} infrared spectra of 217 O-AGB and 98 RSG stars in the Milky Way, Magellanic Clouds and Galactic globular clusters, to explore the onset of crystallinity and investigate how the mineralogy depends on the physical and chemical conditions of the star's envelope. Dust mass-loss rates were established through spectral energy distribution fitting with the {\\sc grams} model grid and the mineralogy of the crystalline features determined from the spectra. The main results of this study are summarised as follows:  \\begin{itemize} \\item We detect crystalline silicates over 3 dex in dust mass-loss rate down rates of $\\sim$10$^{-9}$ M$_\\odot$ yr$^{-1}$. \\item Crystalline silicates are more prevalent in higher mass-loss rates objects, though sources undergoing the highest mass loss do not show the 23-\\mum forsterite feature. This is due to the poor contrast of the low temperature of the forsterite grains above the continuum and in some cases (self-)absorption of the short wavelength ($\\lambda < 25 \\mu$m) crystalline silicate features. \\item The dust mass-loss rate appears to have a greater influence on the crystalline fraction than the gas mass-loss rate. This may suggest that the annealing of amorphous silicate grains by radiation is probably the primary formation mechanism for crystalline silicates in the outflows of AGB and RSG stars. \\item  O-AGB stars have a higher proportion of sources with crystalline silicates features than RSGs, however, there is little variation in the structure of the crystalline silicate dust for O-AGB and RSG stars. \\item We report the presence of a newly detected 22.8 $\\mu $m emission feature in the spectra of Milky Way AGB and RSG stars. \\item We detect a change in the crystalline silicate mineralogy with metallicity, with enstatite seen increasingly at low metallicity, while forsterite becomes depleted. This variation in the olivine-to-pyroxene ratio can be explained by a number of possible mechanisms. \\end{itemize}"
    },
    "1208/1208.1058_arXiv.txt": {
        "abstract": "We report on the first application of the Alcock-Paczynski test to stacked voids in spectroscopic galaxy redshift surveys. We use voids from the Sutter et al. (2012) void catalog, which was derived from the Sloan Digital Sky Survey Data Release 7 main sample and luminous red galaxy catalogs. The construction of that void catalog removes potential shape measurement bias by using a modified version of the {\\tt ZOBOV} algorithm and by removing voids near survey boundaries and masks. We apply the shape-fitting procedure presented in Lavaux \\& Wandelt (2012) to ten void stacks out to redshift $z=0.36$. Combining these measurements, we determine the mean cosmologically induced ``stretch'' of voids in three redshift bins, with $1\\sigma$ errors of 5-15\\%.  The mean stretch is consistent with unity, providing no indication of a distortion induced by peculiar velocities. While the statistical errors are too large to detect the Alcock-Paczynski effect over our limited redshift range, this proof-of-concept analysis defines procedures that can be applied to larger spectroscopic galaxy surveys at higher redshifts to constrain dark energy using the expected statistical isotropy of structures that are minimally affected by uncertainties in galaxy velocity bias. ",
        "introduction": "\\label{sec:introduction}  Characterizing the nature and history of dark energy is perhaps the greatest challenge in the near future of observational cosmology. Many elementary probes now strive to distinguish a cosmological constant from alternative theories of dynamical dark energy or modified gravity. Most probes rely on ``standard candles,'' such as Type Ia supernovae~\\citep[e.g.,][]{AlderingGreg2002}, or ``standard rulers,'' such as radio galaxy diameters~\\citep{Daly2009} or baryon acoustic oscillations (BAO)~\\citep[e.g.,][]{Eisenstein2005, Blake2011, Beutler2011, Anderson2012, MehtaKushalT.2012}. Reviews of dark energy probes, current constraints and forecasts for future experiments include ~\\citet{LinderEricV.2003}, ~\\citet{AlbrechtAndreas2006}, ~\\citet{Frieman2008}, and~\\citet{Weinberg2012}.  Over 30 years ago Alcock \\& Paczynski (1979, hereafter AP) proposed an elegant alternative approach based on a hypothetical population of idealized spheres. Their key insight was that since galaxy spatial positions are inferred from both their angular positions \\emph{and} redshifts, these spheres will appear anisotropic if one adopts an incorrect spacetime metric. Specifically, because line-of-sight distances scale with the inverse Hubble parameter $H^{-1}(z)$ and transverse distances scale with the angular diameter distance $D_A(z)$, their ratio, or \\emph{stretch}, measures the value of the product $H(z)D_A(z)$. In practice, the AP test requires only \\emph{statistical} isotropy of the observed structures, so the test can be implemented with measures of quasar, galaxy, or Ly$\\alpha$ forest clustering, or features in the redshifted 21~cm spectrum~\\citep[e.g.,][]{Hui1999, McDonald1999, EriksenK.A.2005, Nusser2005, Kim2007a, Blake2011, Reid2012}. In this paper we apply the AP test to the~\\citet{Sutter2012} catalog of voids in the galaxy redshift surveys of the Sloan Digital Sky Survey (SDSS;~\\citealt{York2000}).  To date, most applications of the AP test have focused on the autocorrelation function or power spectrum ~\\citep[e.g.,][]{Ballinger1996,Matsubara1996,Matsubara2004}. Specifically, the clearest detections of the AP effect have been found in the two-point correlations of galaxies in the WiggleZ survey~\\citep{Blake2011} and the Baryon Oscillation Spectroscopic Survey~\\citep{Reid2012}. A successful application of the AP test requires handling the large systematic uncertainties caused by peculiar motions, which introduce redshift-space anisotropy that must be disentangled from the AP effect itself. Uncertainties in the peculiar velocity corrections limit the ~\\citet{Blake2011} and~\\citet{Reid2012} studies to large, quasi-linear scales, where the statistical uncertainties are relatively large. An attempt was made recently to apply the AP test to close galaxy pairs~\\citep{Marinoni2010}, but as~\\citet{Bueno2012} point out this method has serious shortcoming due to dynamics at small scales. Additionally, the analysis of~\\citet{Jennings2012} indicates that this method provides relatively weak constraints.  Cosmic voids provide an attractive alternative for applying the AP test, as first proposed by~\\citet{Ryden1995} and discussed extensively by~\\citet{LavauxGuilhem2011}. Voids are the large, underdense regions that occupy a large fraction of the volume of the Universe and are a natural consequence of the hierarchical growth of structure~\\citep{Hausman1983,Thompson2011}. While peculiar velocities modestly affect void shapes~\\citep{Ryden1996,Maeda2011,LavauxGuilhem2011}, voids avoid the regions of high velocity dispersion that have such a large impact on the redshift-space correlation function and power spectrum. Indeed, modeling of peculiar velocities in voids is particularly straightforward since they are still in the quasi-linear regime. In addition, the scale of voids is fairly small, with typical comoving radii $\\sim 10$~\\hmpc~in a densely sampled survey, and they have a large filling factor (i.e., they occupy a majority of the volume of the Universe), amplifying their statistical power relative to other techniques. We can therefore measure the mean void shape with high precision in a large volume survey. ~\\citet{LavauxGuilhem2011} showed that a statistics-limited void AP test can dramatically improve the dark energy constraints from the redshift survey planned for the Euclid satellite~\\citep{Laureijs2011}; the AP test outperforms the BAO constraints from the same survey, even though BAO constraints leverage a known standard ruler, because the scale of the voids is so much smaller than the BAO scale, yielding correspondingly more precise measurements.  We take the void sample for this analysis from the catalog described in~\\citep{Sutter2012}. That work constructed a void catalog from the main galaxy redshift survey ~\\citep{Strauss2002} and the luminous red galaxy (LRG) redshift survey \\citep{Eisenstein2001} of the SDSS Seventh Data Release (DR7;~\\citealt{Abazajian2009}). We identify voids using a modified version of the Voronoi-based {\\tt ZOBOV} algorithm~\\citep{Neyrinck2008}. To compensate for the significant Poisson sampling noise in shape measurements of individual voids~\\citep{Shoji}, we instead measure the mean void shape by ``stacking'' the galaxy distributions of our identified voids in bins of redshift and radius. The SDSS LRG survey is sparse, so at $z \\ga 0.2$ we can only identify large voids, which are limited in number. The combination of moderate redshift leverage and limited statistics prevents us from making a secure detection of the AP effect in this sample, but our proof-of-concept analysis addresses many of the practical issues that will also arise in future data sets at higher redshift.  In the following section we give a brief overview of our method for measuring void shapes and applying the AP test. We review the properties of the~\\citet{Sutter2012} void catalog in Section~\\ref{sec:catalogs}, followed by a presentation of the stacked voids in Section~\\ref{sec:stacks}. We estimate the uncertainty in the stretch measurement and present the AP test as applied to our void stacks in Section~\\ref{sec:ellipticity}. Finally, we offer concluding remarks and a brief discussion of prospects for future surveys in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have performed the first application of the Alcock-Paczynski test to stacked voids to observational data. We applied the AP test by measuring ellipticities of stacked voids using the void catalog of~\\citet{Sutter2012}. The stacking procedure greatly reduces the effects of Poisson noise and allows us to reliably apply the shape-fitting algorithm of~\\citet{LavauxGuilhem2011}. By grouping voids into multiple radial bins we obtain many independent measurements, and by dividing the void catalog into redshift bins we obtain measurements across the full range of the SDSS DR7 main sample and most of the LRG catalog. However, the limited number of voids and the considerable scatter that remains does not allow us to positively identify the AP effect over the redshift range probed by these data.  The SDSS-III BOSS survey~\\citep{Dawson2012} should be a much more powerful basis for void-based AP measurements than the DR7 redshift survey analyzed here.  First, in the range of redshift overlap with our {\\it lrgdim} sample ($z=0.16-0.36$), the space density of BOSS galaxies is a factor $\\sim 3$ higher, which enables identification of smaller (and more numerous) voids and more accurate measurement of void density distributions.  Second, BOSS extends this higher sampling density out to $z\\approx 0.65$, probing a larger comoving volume (and hence more voids) and reaching redshifts where the predicted AP signal is larger, with a stretch factor $e_V(z=0.65) = 1.2$ for a flat-$\\Lambda$ universe with $\\Omega_{\\rm M}=0.27$.  Future ground-based surveys like BigBOSS~\\citep{Schlegel2011} could extend these studies to $z\\approx 1$, while the space-based emission line redshift surveys of Euclid~\\citep{Laureijs2011} and WFIRST \\citep{Green2011} will probe much larger comoving volumes at $z=1-2.5$.  The Fisher matrix analysis of~\\citet{LavauxGuilhem2011} implies that a void-AP analysis of the Euclid survey should yield significantly tighter dark energy constraints than the BAO analysis of the same survey, with a factor of ten improvement in the Dark Energy Task Force~\\citep{AlbrechtAndreas2006} Figure of Merit.  Our initial foray into observational application of this approach highlights two important directions for future investigation. The first is a more detailed study of measurement and parameter fitting procedures and error estimation techniques. Our fitting methods are closely modeled on those of~\\citet{LavauxGuilhem2011}, but there may be other approaches that make better use of the available information, such as using an empirical radial profile in place of our adopted parametric model, changing the radial range of the fit, or downweighting the fluctuations arising from clustered galaxies at the void boundaries.  Alternatively, one could avoid profile fitting entirely and instead use anisotropy of the ``void-galaxy cross-correlation function,'' analogous to the cluster-galaxy cross-correlation but centered at density minima instead of density maxima.  The second direction is a more detailed study of peculiar velocity effects on mean void shapes, examining its dependence on void size and on the spatial and velocity bias of galaxy tracers.  The likelihood analysis in Figure~\\ref{fig:1dlikelihood} allows for an overall velocity distortion factor and therefore effectively uses just the redshift dependence of the signal in Figure~\\ref{fig:hubble} to constrain cosmology, which was appropriate given the systematic effects noticed in~\\citep{LavauxGuilhem2011}. The goal for future analyses should be to apply a theoretically computed velocity distortion correction to each void sample in each redshift bin and marginalize only over the uncertainty in this correction, getting an absolute constraint on the average void stretch, and hence $H(z)D_A(z)$, at each redshift. In the context of halo occupation distribution (HOD) models (e.g.,~\\citealt{Zehavi2011}), we expect galaxies to have the same mean velocities as their parent halos on average, but the velocity dispersion of galaxies could differ from that of the dark matter. This velocity dispersion bias can itself be constrained by redshift-space galaxy clustering~\\citep{Tinker2006}, so we expect the residual uncertainty in peculiar velocity corrections to void shapes to be small, though it may still be the limiting systematic in void-based AP analysis.  The statistical errors of this approach are limited only by the size and redshift range of spectroscopic galaxy surveys, which are expected to grow dramatically in the coming years. Cosmic voids are the converse of galaxy clusters; primordial density minima expand and deepen to form non-linear structures that fill much of the universe and are, in a sense, the most ``dark energy dominated'' regions of the cosmos. The mean shapes of these regions may ultimately provide powerful clues to the nature of the dark energy that pervades them."
    },
    "1208/1208.6352_arXiv.txt": {
        "abstract": "{We develop a general framework for quantifying the many different contributions to the noise budget of an image made with an array of dishes or aperture array stations. Each noise contribution to the visibility data is associated with a relevant correlation timescale and frequency bandwidth so that the net impact on a complete observation can be assessed when a particular effect is not captured in the instrumental calibration. All quantities are parameterised as function of observing frequency and the visibility baseline length. We apply the resulting noise budget analysis to a wide range of existing and planned telescope systems that will operate between about 100 MHz and 5 GHz to ascertain the magnitude of the calibration challenges that they must overcome to achieve thermal noise limited performance. We conclude that calibration challenges are increased in several respects by small dimensions of the dishes or aperture array stations. It will be more challenging to achieve thermal noise limited performance using 15~m class dishes rather than the 25~m dishes of current arrays. Some of the performance risks are mitigated by the deployment of phased array feeds and more with the choice of an $(alt,az,pol)$ mount, although a larger dish diameter offers the best prospects for risk mitigation. Many improvements to imaging performance can be anticipated at the expense of greater complexity in calibration algorithms. However, a fundamental limitation is ultimately imposed by an insufficient number of data constraints relative to calibration variables. The upcoming aperture array systems will be operating in a regime that has never previously been addressed, where a wide range of effects are expected to exceed the thermal noise by two to three orders of magnitude. Achieving routine thermal noise limited imaging performance with these systems presents an extreme challenge. The magnitude of that challenge is inversely related to the aperture array station diameter. } ",
        "introduction": "} Dynamic range limitations in synthesis imaging arise from three distinct categories of circumstance: 1. Instrumental artefacts. 2. Imaging artefacts. 3. Incomplete calibration of the instrumental response. The first category can be addressed by insuring a linear system response to signal levels together with other design measures within the receiver and correlator systems that minimise spurious responses. While challenging to achieve, the engineering requirements in this realm are moderately well defined and this class of circumstance will not be considered further in the current discussion. Some of the specific relevant effects, such as ``closure errors'' and quantisation corrections are discussed in \\citetads{1999ASPC..180..275P}. The second category includes a range of effects, from the well understood smearing effects that result from a finite time and frequency sampling to the imaging challenges associated with non-coplanar baselines as discussed in \\citetads{2008ISTSP...2..647C}. The third category is one that is less well understood and documented. Several important effects in this area are also discussed by \\citetads{1999ASPC..180..275P}; including imperfect polarisation calibration, inadequate (u,v) coverage, and numerical modelling errors. In this study we will attempt to identify and quantify the phenomena that influence synthesis image dynamic range more generally.  The approach we adopt is the consideration of a wide range of calibration issues that influence interferometeric imaging. For each of these effects we develop a parameterised model that quantifies the fluctuation level due to that effect on the measured visibilities, together with its correlation timescale and frequency bandwidth. We then assess the equivalent image noise due to each effect. While some effects contribute directly to the image noise, many contribute to image noise in an indirect manner via the self-calibration process. We demonstrate that both direct (in-field) and indirect (out-of-field) noise contributions typically have a comparable image magnitude if it has proven necessary to employ self-calibration. By tracking each calibration effect individually, it becomes possible to quantify how each contributes to a final image noise level. The purpose of calibration and imaging strategies is to accurately model all of the effects that might limit imaging performance. We do not attempt to evaluate how well any specific algorithm or strategy performs in this regard. We only provide an indication of the precision with which each of these calibration effects must be addressed for an assumed telescope performance specification, so as not to form an ultimate limitation on performance. ",
        "conclusions": "}  \\subsection{Dishes\\label{section:dish}}  Consideration of the various telescope systems in Section~\\ref{sec:exam} allows some general conclusions to be drawn. The ``traditional'' 25~m diameter class dishes of existing synthesis radio telescopes operate in a regime where the thermal noise of the receiver system dominates the noise budget on a self-cal solution interval, typically by a large margin. Furthermore, at least for relatively compact configurations, there is adequate signal-to-noise provided by sources occurring randomly within the field to allow self-calibration to be undertaken at close to the ``natural'' solution interval that keeps time and bandwidth smearing effects at a small fraction of the synthesied beamwidth over the entire primary beam field of view. This constitutes the most advantageous of circumstances, since simple imaging strategies can be employed that use only a single version of the data, in contrast to more complex strategies that might require simultaneous processing of multiple phase centres. Even so, full track continuum observations must overcome several potential obstacles, at the 10 dB level, before thermal noise limited performance is achieved below 2~GHz. This is consistent with practical experience with the JVLA, the WSRT and the ATCA where considerable care must be excercised with precise modelling of sources that occur within the field and in the case of the JVLA, explicit modelling of the non-axisymmetric beam shape and its rotation on the sky during tracking, to keep these effects from becoming dominant. Above about 2~GHz and in more extended array configurations it becomes necessary to introduce more data averaging prior to self-calibration than is optimum for simple imaging strategies. However, only a limited number of complications need be addressed to achieve thermal noise limited performance in full-track continuum observations in these circumstances. The full-track spectral line performance of such arrays is less subject to complications and only requires self-calibration to reach the thermal noise floor below about 2 GHz. Only in the most compact configurations and at frequencies approaching 1~GHz does the Sun begin to present some problems.  This is in stark contrast to the situation encountered with synthesis arrays that employ significantly smaller dishes. The most striking comparison is with the ATA, where the combination of small dish diameter and compact array configuration yield a very different regime. For a nominal self-cal solution interval, the thermal noise is dominated by the far sidelobe response to the entire sky. The same is true even after integrating over an entire 12 hour track, both for spectral line and continuum applications. Continuum observations are further challenged by a wide range of other noise contributions which dominate over the thermal noise. Its not clear whether high dynamic range imaging in this regime will be possible.  The case of 12 to 15~m class dishes in more extended configurations is improved over that with only 6~m dishes. Even so, more than nominal data averaging is required to provide the signal-to-noise needed for self-calibration with this dish size, which will add complexity to wide-field imaging. Higher source modelling precision than what has been needed for the 25~m VLA dishes, by about a factor of two will be required, together with explicit modelling of any main beam asymmetries in cases where the parallactic angle is not fixed on the sky during an observation. In this respect, the ASKAP telescope benefits from both the greater circular symmetry of polarisation beams possible with a phased array feed as well as the polarisation rotation axis that keeps the beam orientation fixed on the sky during tracking. Single pixel fed systems with an $(alt,az)$ mount will face significantly greater challenges to reach the thermal noise floor. Deep field continuum observations will need to achieve extremely high source modelling precision that extends out into the near-in sidelobe region.  Exactly the same considerations apply to the SKA, if 15~m class dishes are deployed as envisaged within the current baseline design. Achieving thermal noise limited continuum performance will be more challenging than with the 25~m dishes of the JVLA and sub-$\\mu$Jy imaging would need to overcome a whole range of complex modelling issues. One of the greatest threats to thermal noise-limited performance is the departure from circular symmetry of the polarisation beams and the rotation of this pattern on the sky. The use of phased array feeds provides one means of reducing the risk associated with beam asymmetries, while the adoption of an equatorial or $(alt,az,pol)$ mount system provides another for eliminating their time variability.  Adoption of a larger dish diameter would provide a significant risk reduction and this option should be considered very seriously during cost/performance optimsation of the final SKA dish design.  \\subsection{Aperture arrays\\label{section:aas}}  The aperture arrays considered in Section~\\ref{sec:exam} display many of the same phenomena. The small station sizes of LOFAR (30.8~m) and particularly the MWA (4.4~m) result in visibilities that are completely dominated by noise-like modulations from sources appearing in the far sidelobe response. Very extensive modelling of the visibility response to these distant sources becomes necessary before self-calibration within the main beam can be contemplated. A major additional complication is introduced by the geometric foreshortening of each station aperture while tracking a source. While an absence of moving parts is often cited as a major advantage of these systems, the continuously changing response of the sky-telescope system is the price that is paid. Efforts are underway to improve predictive modelling of both the far sidelobe response patterns \\citepads{5355494} as well as compensation for the time-variable main beam shape \\citepads{2009IEEEP..97.1472R}. Current LOFAR performance employing baselines up to 30~km and a relative bandwidth, $\\Delta\\nu/\\nu \\sim 0.35$, in a well-characterised reference field have reached noise levels of about 100--180 $\\mu$Jy/beam, or within a factor of about 1.4 of the thermal noise \\citepads{2013arXiv1301.1630Y} for a 10 hour effective duration observation.  While the calibration challenge is daunting for LOFAR, where some 20 dB of reduction is needed to approach the thermal noise on a solution interval, it is more extreme for the MWA where 30-40 dB are needed. Although self-calibration seems to be required to reduce the magnitude of gain calibration errors it does not appear to be a viable option for the MWA. The source population responsible for residual visibility errors is so large, consisting of more than $10^4$ relevant sources above the horizon at any moment, that it may not prove practical to model the time-dependent system response to these sources at all. This may be an intrinsically under-determined problem with the finite information available within the solution interval over which the instrumental response is time-invariant. While it has been suggested \\citepads{2009MNRAS.400.1814M} that alternate calibration and imaging strategies might be utilised to circumvent these challenges, the fundamental question of sufficient independent data constraints on a solution timescale has not been demonstrated in either theory or practise.  The SKA1-Low concept calls for much larger aperture arrays, of 180~m diameter, to combat the problem of far sidelobe response. While this approach helps, it is still the case that many noise contributions exceed the thermal noise, both on a self-cal solution interval as well as in full-track spectral line and continuum images. Some 30 dB of attenuation is necessary for a wide range of phenomena to approach the thermal noise in a single 4 hour track. This will be extremely challenging to overcome, since even the 10 dB level issues faced by the VLA 25~m dishes at GHz frequencies have required many years of algorithm development to partially address. While large station diameters seem to offer the best prospects, it will still require very significant progress in both algorithm development and efficient computational implementations over what is available currently to approach thermal noise limited performance for $\\nu < 200$ MHz."
    },
    "1208/1208.3741_arXiv.txt": {
        "abstract": "{% We have used two robotic telescopes to obtain time-series high-resolution optical echelle spectroscopy and $VI$ and/or $by$ photometry for a sample of 60 active stars, mostly binaries. Orbital solutions are presented for 26 double-lined systems and for 19 single-lined systems, seven of them for the first time but all of them with unprecedented phase coverage and accuracy. Eighteen systems turned out to be single stars. The total of 6,609 $R$=55,000 echelle spectra are also used to systematically determine effective temperatures, gravities, metallicities, rotational velocities, lithium abundances and absolute \\Halpha -core fluxes as a function of time. The photometry is used to infer unspotted brightness, $V-I$ and/or $b-y$ colors, spot-induced brightness amplitudes and precise rotation periods. An extra 22 radial-velocity standard stars were monitored throughout the science observations and yield a new barycentric zero point for our STELLA/SES robotic system. Our data are complemented by literature data and are used to determine rotation-temperature-activity relations for active binary components. We also relate lithium abundance to rotation and surface temperature. We find that 74\\% of all known rapidly-rotating active binary stars are synchronized and in circular orbits but 26\\% (61 systems) are rotating asynchronously of which half have $P_{\\rm rot}>P_{\\rm orb}$ and $e>0$. Because rotational synchronization is predicted to occur before orbital circularization active binaries should undergo an extra spin-down besides tidal dissipation. We suspect this to be due to a magnetically channeled wind with its subsequent braking torque. We find a steep increase of rotation period with decreasing effective temperature for active stars, $P_{\\rm rot} \\propto T_{\\rm eff}^{-7}$, for both single and binaries, main sequence and evolved. For inactive, single giants with $P_{\\rm rot}>100$~d, the relation is much weaker, $P_{\\rm rot} \\propto T_{\\rm eff}^{-1.12}$. Our data also indicate a period-activity relation for \\Halpha\\ of the form $R_{\\rm H\\alpha} \\propto P_{\\rm rot}^{-0.24}$ for binaries and $R_{\\rm H\\alpha} \\propto P_{\\rm rot}^{-0.14}$ for singles. Its power-law difference is possibly significant. Lithium abundances in our (field-star) sample generally increase with effective temperature and are paralleled with an increase of the dispersion. The dispersion for binaries can be 1--2 orders of magnitude larger than for singles, peaking at an absolute spread of 3 orders of magnitude near $T_{\\rm eff}\\approx$5000~K. On average, binaries of comparable effective temperature appear to exhibit 0.25~dex less surface lithium than singles, as expected if the depletion mechanism is rotation dependent. We also find a trend of increased Li abundance with rotational period of form $\\log n (\\mathrm{Li}) \\propto -0.6 \\ \\log P_{\\rm rot}$ but again with a dispersion of as large as 3--4 orders of magnitude.} ",
        "introduction": "In a previous paper (Strassmeier et al.~\\cite{ca2}; paper~I) we reported on radial and rotational velocities, chromospheric emission-line fluxes, lithium abundances, and rotation periods of a total sample of 1,058 G5--K2 dwarfs, subgiants, and giants based on 1,429 moderate-resolution KPNO coud\\'e spectra and 8,038 Str\\\"omgren $y$ photometric data points. The aim of this survey was to detect new candidates for Doppler imaging but, besides the discovery of 170 new variable stars and 36 new spectroscopic binaries, the more intriguing result was that 74\\%\\ of the G-K stars with Ca\\,{\\sc ii} H\\&K emission also showed significant lithium on their surface. However, G-K giants should have very few lithium on their surface because of convective mixing. Theoretical models predict that surface lithium has to be diluted by many factors once a star arrives at the bottom of the red giant branch (Iben \\cite{iben67}, Charbonnel \\& Balachandran \\cite{cha:bal}). Out of the 21 Doppler imaging candidates found, just four stars were single stars, three of them evolved, the rest were spectroscopic binaries, but all four single stars had very strong lithium.  Despite that it is generally acknowledged that higher than normal lithium abundance is common among magnetically active stars, no unique correlation with rotation rate was found after Skumanich's (\\cite{sku}) original discovery. Recently, White et al. (\\cite{white2007}) revisited this issue in their sample of solar-type dwarfs but no such correlation was found. Almost all surveys just revealed trends, if at all, and even these appear to be of different quality (e.g. L\\'ebre et al.~\\cite{lebre2006}; B\\\"ohm-Vitense~\\cite{boehm}; do Nascimento et al.~\\cite{dona2000}, \\cite{dona2003}; De Medeiros et al.~\\cite{dem2000}; De Laverny et al.~\\cite{dela2003}; Randich et al. \\cite{ran:gra}). The comprehensive survey of nearby giants by Luck \\& Heiter (\\cite{luc:hei}) did not even show a trend. However, the line broadening in their stellar sample was just 3--7~\\kms\\ and likely too narrow a range to see a trend. B\\\"ohm-Vitense~(\\cite{boehm}) suggested that the steep decrease of $v\\sin i$ in early G giants as well as in Hyades F dwarfs at effective temperatures cooler than $\\approx$6,450~K, i.e. at their lithium dip at about the same temperature, are the result of deep mixing and related to the merging of the hydrogen and the helium convection zones. More recently, Takeda et al. (\\cite{tak:hon}) announced evidence for a (positive) correlation of Li abundance with rotational velocity in a sample of solar-analog stars.  The spectroscopic survey of 390 solar-like dwarf stars by White et al.~(\\cite{white2007}) included 28 of the stars in our survey. Relating Ca\\,{\\sc ii} H\\&K radiative losses to stellar rotation, White et al.~(\\cite{white2007}) found a saturation of chromospheric emission for rotational velocities above approximately 30~\\kms . In an earlier paper, Strassmeier et al.~(\\cite{str:han}) verified that evolved stars obey qualitatively the same scaling of Ca\\,{\\sc ii}-K-line flux with stellar rotational velocity or period as do main-sequence stars (see, e.g., Mamajek \\& Hillenbrand \\cite{mam:hil}, Pace \\& Pasquini~\\cite{pac:pas}, or Pizzolato et al.~\\cite{piz:mag} for a summary). No qualitative difference was found between single evolved stars and their equally rapidly rotating counterparts in a spectroscopic binary. However, large scatter indicated that rotation might not be the only relevant parameter. Based on a sample of 22 intermediate-mass G and K giants in close binaries, Gondoin (\\cite{gon}) not only verified the rotational dependency of (coronal) X-ray surface flux but also found a dependency on surface gravity. Such a dependence could stem from the effect of gravity on coronal electron density and on the overall sizes of coronal loops.  Massarotti et al. (\\cite{mas:lat}) reported rotational and radial velocities for 761 giants within 100~pc of the Sun. They found that all binaries in their sample with periods less than 20 days have circular orbits while about half the orbits with periods between 20--100 days still showed significant eccentricity. They also found evidence that the rotational velocity of horizontal branch stars is larger than that of first-ascend giants by a few \\kms . Earlier, De~Medeiros et al. (\\cite{dem2002}) presented a study of 134 late-type giants in spectroscopic binaries and found a considerable number of G-K giant stars with moderate to moderately-high rotation rates. These rotators have orbital periods shorter than 250 days and circular or nearly circular orbits and appear to be synchronized with the orbit.  The present paper follows up on the newly identified spectroscopic binaries with active components from our paper~I. Its direct aim is to determine their orbits on the basis of high-precision radial velocities and to separate their component's rotation and activity tracers along with other absolute astrophysical parameters. Only with precise stellar parameters can we directly compare binary components with single stars and then be aware of the spectrum contamination from unknown secondaries or even tertiary stars. We recall that an unknown continuum contribution from a secondary star impacts on the determination of the effective temperature, gravity etc. and could together drastically alter the derived lithium abundances and thereby mask any relation if present. In Sect.~\\ref{S2} we restate our sample selection criteria and give a summary of the target stars. In Sect.~\\ref{S3} we describe the new observations and in Sect.~\\ref{S4} we derive basic quantities from the spectra and the light curves. These include radial velocities, orbital parameters, rotational velocities and photometric periods, stellar atmospheric parameters like temperature, gravity and metallicity, lithium abundances, and absolute \\Halpha -core fluxes. Sect.~\\ref{S5} lists notes to individual stars. Sect.~\\ref{S6} presents the analysis in terms of rotation, temperature, activity, and lithium-abundance relations. Finally, Sect.~\\ref{S7} summarizes our findings and conclusions. ",
        "conclusions": "\\label{S7}  We present and analyzed high-resolution spectra from the STELLA robotic telescope and determined orbital elements for 45 binaries in the period range 0.29--6450~d, many for the first time but all with substantially increased precision and likely also increased accuracy. Additional photometric monitoring of many of these stars with our APTs allowed the determination of precise rotation periods.  We summarize our findings as follows. \\begin{itemize} \\item STELLA/SES achieves a radial-velocity precision for an individual radial-velocity observation of $\\approx$30~\\ms\\ without an iodine cell or a simultaneously recorded comparison spectrum. It enables rms values of as small as $\\approx$40~\\ms\\ for fits of orbital elements. For one of the favorable cases in this paper (the SB2 binary HIP~77210 with $P_{\\rm orb}$=9.9~d) the minimum primary mass is precise to 0.12\\%\\ and the secondary mass to even 0.07\\% . \\item The radial-velocity zero-point difference between STELLA and CORAVEL measurements is +0.503~\\kms . \\item The bulk (74\\%) of rapidly-rotating active stars in binaries are synchronized and in circular orbits. \\item About 26\\%\\ (61 targets) of the binaries in our full sample are asynchronous rotators. About half of them have $P_{\\rm rot}>P_{\\rm orb}$ and of these, all but two have $e>0$. As rotational synchronization is predicted to occur before orbital circularization active binaries must have gone through an extra spin-down phase besides tidal dissipation and we suggest this to be due to a magnetically channeled wind with its subsequent braking torque. Such braking depends on the stellar magnetic-field geometry and would result in very different braking efficiencies for stars of otherwise nearly identical physical parameters. \\item We find a steep increase of $P_{\\rm rot}$ for lower $T_{\\rm eff}$ in both single stars and binary components. This verifies earlier claims for active single stars with known cycles and shows that it is also the case in our active binaries. A functional dependence of $P_{\\rm rot}\\propto T_{\\rm eff}^{-X}$ is suggested, where $X=7.2$ for single stars and $X=7.0$ for binary stars. The difference in $X$ is not significant though. \\item A relation between \\Halpha -core flux and rotational period is evident for both single stars and binaries. Power-law fits suggest $R_{\\rm H\\alpha} \\propto P_{\\rm rot}^{-0.24}$ for binary stars and $R_{\\rm H\\alpha} \\propto P_{\\rm rot}^{-0.14}$ for single stars, where $R$ is \\Halpha -core flux normalized to $\\sigma T_{\\rm eff}^4$. The difference of the exponents is only weakly significant on the 3$\\sigma$ level. \\item Our data suggest that Li abundances of active binary components also increase with effective temperatures, as known for single stars. However, binaries appear to have on average 0.25~dex lower Li abundances than single stars of same effective temperature, which is what we expect if the depletion mechanism is rotationally enhanced in binary stars. The dispersion for binaries appears to peak around $T_{\\rm eff}\\approx$5000~K, amounting to three orders of magnitude. This dispersion decreases towards both ends of the temperature range between 4500~K to 6500~K. Our single-star sample follows this trend as well but with an order of magnitude smaller dispersion. We can not separate age effects in our sample but attribute a significant fraction of the dispersion to it. \\item Rotational dependency of Li abundance in binaries is also evident and suggests a law of form $\\log n (\\mathrm{Li}) \\propto -0.6 \\ \\log P_{\\rm rot}$ as a guideline. However, a dispersion of up to 3--4 orders of magnitude explains why no quantitative relations were found so far. The dispersion appears largest for binaries with periods shorter than $\\approx$10~d and still amounts to 2--3 orders of magnitude for periods larger than $\\approx$10~d. \\end{itemize}"
    },
    "1208/1208.0742_arXiv.txt": {
        "abstract": "During 2010-2011, the Medicina 32-m dish hosted the 7-feed 18-26.5 GHz receiver built for the Sardinia Radio Telescope, with the goal to perform its commissioning. This opportunity was exploited to carry out a pilot survey at 20 GHz over the area for $\\delta >~+72.3^\\circ$. This paper describes all the phases of the observations, as they were performed using new hardware and software facilities. The map-making and source extraction procedures are illustrated. A customised data reduction tool was used during the follow-up phase, which produced a list of 73 confirmed sources down to a flux density of 115 mJy. The resulting catalogue, here presented, is complete above 200 mJy. Source counts are in agreement with those provided by the AT20G survey. This pilot activity paves the way to a larger project, the K-band Northern Wide Survey (KNoWS), whose final aim is to survey the whole Northern Hemisphere down to a flux limit of 50 mJy (5$\\sigma$). ",
        "introduction": "Extragalactic radio sources extracted from high-frequency ($>10$~GHz) surveys are expected to have a major impact on astrophysics. They can provide samples of rare classes of sources with flat or inverted spectrum that, at low frequencies, are swamped by more numerous populations which fade away as the frequency increases (for a review see De Zotti et al. 2010).  They can hence open a window on new classes of sources, such as those with strong synchrotron or free-free self-absorption corresponding to both very early phases of nuclear radio-activity (extreme GHz Peaked Spectrum - GPS - sources or high-frequency peakers) and late phases of the evolution of Active Galactic Nuclei (AGNs), characterized by low accretion/radiative efficiency (ADAF/ADIOS sources), as well as to early phases of the evolution of radio afterglows of gamma-ray bursts. In this context the comparison with the on-going Fermi observations will yield very interesting results (see Mahony et al. 2010).  These sources also play a vital role in the interpretation of temperature and polarisation maps of the Cosmic Microwave Background (CMB). Extragalactic point sources are one of the major foreground emissions (Planck collaboration 2011, Leach et al. 2008, Toffolatti et al. 2005); the knowledge of their positions and flux densities is crucial to remove their contribution and to estimate the residual error due to faint and unresolved components in CMB maps. As the source population composition changes at high frequency, cleaning procedures based on lower frequency catalogues are unreliable, making it essential to carry out surveys at frequencies close to the CMB window (centred at 60-70~GHz). In addition, as the Planck satellite is now active, the realisation of a coeval 20-GHz blind survey helps, when selecting flux density limits, to avoid errors induced by high-frequency variability.  High frequency sky surveys have become feasible very recently. Because of the faint signal, the existing surveys at 10-100~GHz usually cover small areas with good sensitivity (e.g., VSA at 34~GHz with $S_{lim}  =$ 100~mJy, Gawro\\'nski et al. 2010) or consist in all-sky shallow surveys (e.g., WMAP at 23, 33, 41, 64, 94 GHz  with $S_{lim} >1$~Jy, Wright et al. 2009 a,b and ERCSC at 30, 44, 70, 100~GHz - Planck collaboration 2011 a,b,c). The only exception to this is the all-southern-sky Australia Telescope 20 GHz survey (AT20G), which observed the entire southern sky with the Australia Telescope Compact Array, detecting around 6000 sources down to a flux density limit of 50~mJy  (Murphy et al. 2010, Massardi et al. 2011).  This calls for a northern sky survey with equivalent sensitivity to complete the coverage of the entire sky. Such a completeness is particularly important for several aims for which a statistical information is not sufficient, like the study of the SED of peculiar objects, the selection of samples at high radio-frequency for the northern or whole sky. A precise position of all the sources is also required to flag out the contaminated pixels from CMB maps.  The availability of a K-band (namely 18-26.5 GHz) multi-feed receiver installed on a medium-sized antenna as the Medicina 32-m dish, having a beamsize of 1.6~arcmin @~21~GHz, gave us the possibility to execute a pilot survey to verify the receiver performance, together with new software tools, while exploring a sky area which had never been extensively observed at these frequencies. This test activity paved the way to a larger project, the K-band Northern Wide Survey, which aims at performing a blind survey over the whole northern hemisphere, with a sensitivity of 50~mJy (5$\\sigma$).  The outline of the paper is the following. After a description of the system capabilities (\\S \\ref{sec:instrument}), we summarise in \\S \\ref{sec:survey} the survey strategy, including the map-making and source extraction techniques applied to achieve the list of candidate sources. The 20-GHz follow-up observations are described in \\S \\ref{sec:followup}, together with the automatic pipeline produced for the data reduction. Finally, the catalogue is presented in \\S \\ref{sec:catalogue} and results are summarised in \\S \\ref{sec:conclusions}. A separate paper by Ricci et al. (hereafter 'Paper II') illustrates the detailed spectral index analysis of the sources. ",
        "conclusions": "\\label{sec:conclusions} Our project served as a fundamental test for the hardware and software facilities which were being commissioned in Medicina. The survey/mapping initial phase and the subsequent multi-frequency follow-up allowed us to deeply test the main continuum observing modes. They also helped us in improving our knowledge of the influence of RFI over high-frequency acquisitions performed with unprecedented sensitivity, as these single-dish facilities were not previously available in Medicina. Map-making, source extraction and data reduction tools, both known and custom-developed, were tested and debugged, allowing us to fine tune the whole processing phase going from data acquisition to the reduction and calibration of high-frequency maps and cross-scans (including the crucial issues related to atmospheric opacity). The survey, performed exploiting an ad hoc observing technique, produced a shallow map of the region for $\\delta > 72.3 ^\\circ$, with an average noise level of 15.4~mJy. A selection of the extracted source candidates was followed-up, leading to the confirmation of 73 sources, down to a flux density of 115~mJy. The characteristics of the map, together with a strict selection of the source candidates imposed by the short commissioning time available, produced this final catalogue of bright sources, which is complete for $S_{lim} = 200$~mJy. It must be noticed that this region of the sky had never been extensively observed at 20~GHz down to this flux density limit, thus this catalogue constitutes a useful reference for spectral studies of the listed sources. A separate paper (Ricci et al., in prep.) illustrates the multi-frequency observations and spectral analysis carried out for these sources."
    },
    "1208/1208.5484_arXiv.txt": {
        "abstract": "Effective supersymmetry(SUSY) where stop is the lightest squark may run into a two-loop tachyonic problem in some $Z'$ mediation models. In addition, a large $A$ term or/and a large stop mass are needed to have $\\sim$ 126 GeV Higgs boson with three families of quarks and leptons. Thus, we suggest an inverted effective SUSY(IeffSUSY) where stop mass is larger compared to those of the first two families. In this case, it is possible to have a significant correction to the anomalous magnetic moment of muon. A three family IeffSUSY in a $Z'$ mediation scenario is explicitly studied with the $Z'$ quantum number related to $B-L$. Here, we adopt both the $Z'$ mediation and gravity mediation where the $Z'$ mediation is the dominant one for stop, while the gravity mediation is the dominant one for the muonic leptons and Higgs multiplets. We present a numerical study based on a specific anomaly free model, and show the existence of the parameter region where all the phenomenological conditions are satisfied. ",
        "introduction": "The recent LHC reports hint the Higgs boson mass at 125--127 GeV \\cite{CERNJuly4}. This small Higgs boson mass compared to the Planck mass needs a huge hierarchy of mass scales, inviting solutions of the hierarchy problem. Supersymmetry(SUSY) has been considered to be the most attractive one among the hierarchy solutions, but the LHC data is not consistent with the constrained minimal supersymmetric standard model(CMSSM) prediction in the region $ \\Mgl\\msq \\lesssim 1\\tev^2$. A small Higgs boson mass ($m_h\\simeq 0.126\\,\\tev$) needs a large stop mass or/and a large $A$-term in the CMSSM.  The LHC hints toward large squark masses are usually interpreted as a large mass limit for the first family squarks. The third family squarks have much lower exclusion bound \\cite{LHC3rd} than those of the first two families. The current bound for $m_{\\tilde q_{1,2}}$ is usually taken as 1.5\\,TeV \\cite{LHCsquark}. Thus, the squark masses of the third family can be below 1\\,TeV in principle, which has been proposed long time ago as the effective SUSY(effSUSY) \\cite{effSUSY95}.  So, if Nature has low energy SUSY as a solution to the hierarchy problem, the previously considered attractive models have been the effSUSY where only the third family squarks are in the reach of the LHC search.  Another LHC hint is the possibility of light Higgs boson whose mass is around 126 GeV \\cite{CERNJuly4}. However, a light Higgs boson as heavy as 126 GeV is difficult to obtain in the CMSSM, mainly due to the tree level mass bound $m_h^{\\rm tree}\\le M_Z$. The loop corrections raise the Higgs mass but a fine-tuning is needed to raise it above 120 GeV \\cite{HiggsBound}. This has led to scenarios with a large $A$-term from the large top Yukawa coupling and/or a large stop mass. The effSUSY through gauge mediation cannot lead to a large $A$-term since the gravity mediation for the squark mass generation is assumed to be sub-dominant compared to that of the gauge mediation. Also, the effSUSY assumes a relatively small stop mass. If gauge mediation is effectively achieved by a family-dependent $Z'$ mediation \\cite{Jeong11}, then there is another problem that most scalar particles become tachyons, if  the  two-loop contributions are included.  A large hierarchy of soft masses between different families is easily realized by imposing family dependent $Z'$ charges in the $Z'$ mediation  \\cite{Lang08,Zpfamind}. If the soft masses of some specific family are much smaller than those of another family, the heavy soft mass term contributes to the light soft masses at the two loop level through the SM gauge group interaction \\cite{ArkaMura97}. Therefore, if the light scalars are not charged under U(1)$_{Z'}$ in the $Z'$ mediation scenario, they become tachyonic when two loop effects are taken into account. This two-loop tachyonic problem may not be present in the $Z'$ D-term breaking \\cite{Hisano99}. However, the model building along the D-term breaking may be more complicated than the method we introduce below with three chiral families of quarks and leptons.  The family-dependent $Z'$ mediation is so easily realized in string models \\cite{Jeong11} that we consider its realization a {\\it natural} one. As mentioned above, however, we need a large $A$-term and/or large stop masses from the LHC constraints.  Here, we implement the large $A$-term by the gravity mediation. Both for the $Z'$ mediation and the gravity mediation, the same dynamical SUSY breaking scale applies, which is assumed to be around $\\Lambda_h\\simeq 10^{13}\\,\\gev$ \\cite{Nilles84}. Gravity mediation with this dynamical SUSY breaking scale sets the scale for the $A$-term. To radiatively raise the Higgs boson mass sufficiently above $M_Z$, we need large stop masses. So, the family-dependent $Z'$ mediation is of the form `inverted', in the sense that the 3rd family squarks are heavy compared to the first two family squarks. We call this scenario {\\it inverted effective SUSY}(IeffSUSY). If stops are the heaviest sfermions, we use the term IeffSUSY irrespective of the order of the remaining sfermion masses.\\footnote{Here, `inverted' is used just for the heaviest stops since effSUSY has been used for the lightest stops among sfermions \\cite{effSUSY95}.}  In the family-dependent $Z'$ mediation scenario, the 3rd family members can be made heavy by assigning large $Z'$ quantum numbers to them while keeping the first two family members to carry very small $Z'$ quantum numbers, which is the opposite view taken from that of Ref. \\cite{Jeong11}. % So, the first two family squarks obtain masses predominantly via the gravity mediation. [Note that the 3rd family members get the additional contribution through the $Z'$ mediation.] For the gravity mediation, the messenger scale is considered to be the Planck mass $M_P=2.44\\times 10^{18}\\,\\gev$. For the  $Z'$ mediation, the messenger scale is another parameter $\\Mmess$. If $\\Mmess\\lesssim \\frac{1}{10}M_P$, we can achieve a reasonable IeffSUSY. The messenger scales and the visible sector masses are depicted schematically in Fig. \\ref{fig:Mmess}, where the visible and the hidden sectors do not communicate directly as emphasized by the thick brown bar in Fig.  \\ref{fig:Mmess}.  Assuming that the lowest messenger scale $\\Mmess$ is significantly separated from the other messenger scales, the low energy spectra is dominated by the scale $\\Mmess$ of the $Z'$ mediation. In this sense, we argue that the $Z'$ mediation arises {\\it naturally} from an ultraviolet completed theory, as far as the lowest messenger scale $\\Mmess$ is sufficiently separated from the other messenger scales.  Now, the two-loop tachyons are made stable by the positive soft mass arising from the gravity mediation. In addition, in this IeffSUSY the muon $g-2$ deviation \\cite{gmtwoBNL} from the SM estimation can be made significant through the light gaugino masses, which arise at the two-loop in the $Z'$ mediation \\cite{Lang08}, and the light smuon ($\\tilde\\mu$) and scalar-muonneutrino ($\\tilde\\nu_2$) masses.  \\begin{figure}[!t] \\begin{center} \\begin{tabular}{c} {\\includegraphics[width=0.35\\textwidth]{fig1MessScale.eps}} \\end{tabular} \\end{center} \\caption{ {A schematic view of messenger scales and the IeffSUSY spectra. The lowest scale, assumed to be separated somewhat from the others, is called $\\Mmess$. } }\\label{fig:Mmess} \\end{figure} ",
        "conclusions": "\\label{sec:Conclusion}  In view of the observed 126 GeV Higgs boson \\cite{CERNJuly4} which is relatively heavy in the SUSY scenario, we introduced a heavy stop scheme which is the opposite view from the popular effSUSY idea. In the $Z'$ mediation scenario, we achieve this IeffSUSY explicitly with the U(1)$_{Z'}$ quantum numbers shown in Table \\ref{table:BmL}. With the quantum numbers of Table \\ref{table:BmL}, it is possible to have relatively light smuons ($\\sim 2-3\\,\\tev$) and neutralino ($\\sim 1.2\\,\\tev$), and hence can find a parameter region where a significant correction to the anomalous magnetic moment of muon can result."
    },
    "1208/1208.0604_arXiv.txt": {
        "abstract": "We present a full Bayesian algorithm designed to perform automated searches of the parameter space of caustic-crossing binary-lens microlensing events. This builds on previous work implementing priors derived from Galactic models and geometrical considerations. The geometrical structure of the priors divides the parameter space into well-defined boxes that we explore with multiple Monte Carlo Markov Chains. We outline our Bayesian framework and test our automated search scheme using two data sets: a synthetic lightcurve, and the observations of OGLE-2007-BLG-472 that we analysed in previous work. For the synthetic data, we recover the input parameters. For OGLE-2007-BLG-472 we find that while $\\chi^2$ is minimised for a planetary mass-ratio model with extremely long timescale, the introduction of priors and minimisation of BIC, rather than $\\chi^2$, favours a more plausible lens model, a binary star with components of $0.78$ and 0.11~$\\msun$ at a distance of $6.3$ kpc, compared to our previous result of $1.50$ and $0.12~\\msun$ at a distance of 1~kpc. ",
        "introduction": "\\label{sec:intro}  Gravitational microlensing \\citep{einstein36} is a well-established technique to detect extrasolar planets (e.g. \\citealt{maopaczynski91}, \\citealt{beaulieu06}, \\citealt{muraki11}), and is complementary to other methods, being able to probe low-mass cool planets that are inaccessible to them from the ground. This allows us to carry out statistical studies of planets of all masses located at a few AU from their host star \\citep{cassan12}. Microlensing occurs when one or several compact objects are located between a source star and the observer, leading to a gravitational deflection of the light from the source star by the ``lens\" objects. As the source and lens move in and out of alignment, this deflection is observable in the form of a simple characteristic brightening and fading pattern when the lensing object is a single star (\\citealt{paczynski86}), but takes a much more complex form when the lens is made up of more than one object. When that happens, the lightcurve typically features ``anomalies\", which can be modelled to determine the nature of the lensing system. One of the configurations that can lead to anomalies is when the lensing system contains one or more planets. In order to determine the properties of these planets, the anomalies must be analysed through detailed modelling; this paper is concerned with cases where the lens consists of two components.  Analysing anomalous microlensing lightcurves can be a significant computational challenge for a number of reasons. The calculation of a full binary-lens lightcurve, including the effects of having an extended source, is an expensive process computationally, and the parameter space to be explored is complex, with several degeneracies (e.g. \\citealt{kubas05}). This is the case even when second-order effects, such as that of parallax due to the Earth's orbit or orbital motion in the lensing system, are ignored.  A significant number of the $\\sim 1500$ microlensing events now being discovered by survey teams in a season exhibit anomalies due to stellar or planetary companions to the lens star. Many of these are caustic-crossing events in which the lightcurve exhibits rapid jumps, brightening when a new pair of images forms and fading when two images merge and disappear. \\cite{cassan08} introduced an advantageous parameterisation for caustic-crossing events by linking two parameters, $\\tin$ and $\\tout$, to the caustic-crossing times and two parameters, $\\sin$ and $\\sout$, to the ingress and egress points where the source-lens trajectory crosses the caustic curve. These parameters make it easier to locate all possible source-lens trajectories that fit the observed caustic-crossing features.  \\cite{kains09} used the \\cite{cassan08} parameters to analyse the observed lightcurve of the microlensing event OGLE-2007-BLG-472, which exhibits two strong caustic-crossing features separated by about 3 days. This short duration suggested that the anomaly could be due to the source crossing a small planetary caustic, motivating detailed modelling to rule out alternative binary star lens models. The lowest-$\\chi^2$ model has a planetary mass ratio, but an extremely long event timescale, $\\te \\sim 2000$ days, much longer than the 2-200~day range typical of Galactic Bulge microlensing events. On this basis \\cite{kains09} rejected the global $\\chi^2$ minimum by placing an ad-hoc 300~day cutoff on $\\te$, and suggested that a Bayesian approach including appropriate priors on all the parameters would more naturally shift the posterior probability to local $\\chi^2$ minima with less extreme parameters.  \\cite{cassan09} derived analytic formulae for the prior $\\priorp{\\sin,\\sout}$ corresponding to a uniform and isotropic distribution of lens-source trajectories, which are specified by an angle $\\alpha$ and impact parameter $\\uz$. A suitable prior on $\\te$ arises by using a model of microlensing in the Galaxy to determine distributions for the lens and source distances and their relative proper motion, or alternatively by using a parameterised model fitted to the observed distribution of $\\te$ among all the events found in the microlensing survey. In either case a prior on $\\te$ effectively penalises very long and very short events, lowering the posterior probability of the $\\te\\sim2000$~d global $\\chi^2$ minimum found for OGLE-2007-BLG-472 and favouring local minima with more typical event timescales. Priors on other parameters can also be derived from models of stellar population synthesis such as the Besan{\\c c}on model \\citep{robin03}, which we use in this work.  In this paper we develop further the Bayesian analysis of caustic-crossing events, exploiting intrinsic features of the $\\priorp{\\sin,\\sout}$ prior to specify and test a procedure suitable for automatic exploration of the full parameter space. We test the procedure using synthetic lightcurve data, and we re-analyse the OGLE-2007-BLG-472 data to compare the results of maximum likelihood analysis ($\\chi^2$ minimisation) with the full Bayesian analysis including appropriate priors. ",
        "conclusions": "\\label{sec:conclusion}  The modelling results for the two datasets presented here indicate that our algorithm is successful in locating minima throughout the parameter space, and the subdivision of the prior maps ensures that all possible source trajectories through the caustics are explored. Furthermore, the use of Bayesian priors allows us to incorporate information on the event timescale distribution, as well geometrical information on the concavity of caustics.  Although the sampling rate for our synthetic lightcurve data is not particularly high compared to what can now be achieved by survey and follow-up teams, our algorithm located a well-defined minimum near the true minimum, with a grid search of the $(d,q)$ parameter space and MCMC runs to sample the posterior probability in the region around each local minimum.  In our re-analysis of the OGLE-2007-BLG-472 data, we improve upon the posterior map calculated in \\cite{kains09} for OGLE-2007-BLG-472 because we now use an MCMC run for each prior sub-box separately rather than just a single one per $(d,q)$ grid point. We find that changing the badness-of-fit statistic leads to important changes in the posterior $P(d,q|D)$ maps. In particular, the model with lowest $\\chi^2$ has a planetary mass ratio and an implausibly long $\\te\\sim2000$~d. Adding priors dramatically shifts the location of the best-fit model, lowering the timescale to $\\te\\sim70$~d. Using a Bayesian approach to penalise models with improbable parameters leads to best-fit parameters corresponding a binary star lens with $0.78$ and 0.12~$\\msun$ components at a distance of $\\sim5.9$~kpc, and a more typical event timescale $\\te\\sim 70$~d. The only remarkable parameter is a rather high blending fraction, which could arise from either the lens itself or a closely blended third star. The new model is very different from that found by \\cite{kains09}, which characterised the lens as a binary star with components masses of $1.50$ and $0.12 \\msun$ at a distance of 1 kpc.  The development of automated algorithms for real-time modelling such as that presented here allows observers to receive feedback on ongoing anomalous microlensing events, and ensure that important features predicted by real-time modelling are not missed. This makes it much easier to assess the nature of the lensing system more rapidly and allocate observing time to targets more effectively. When observational coverage is not complete, or when the  $\\chi^2$ alone is not sufficient as a criterion for badness-of-fit, statistics like the those we use in this paper could help to assess reliably alternative models. Furthermore, provided that the chosen priors are appropriate, comparing the resulting posterior maps of using different statistics allows for a useful test of a given model's robustness."
    },
    "1208/1208.1508_arXiv.txt": {
        "abstract": "{ Current atmospheric and evolutionary models for low-mass stars and brown dwarfs rely on approximate assumptions on the physics of the stellar structure and the atmospheric radiative transfer. This leads to biased theoretical predictions on the photospheric Spectral Energy Distributions of these system, especially when applied to low surface gravity objects such as Pre-Main Sequence (PMS) stars, and affects the derivation of stellar parameters from photometric data. } { Our main goal is to correct the biases present in the theoretical predictions for the near-IR photometry of low-mass PMS stars. Using empirical intrinsic IR colors, we assess the accuracy of current synthetic spectral libraries and evolutionary models. We investigate how the uncertainty in the intrinsic colors associated with different PMS models affect the derivation of the Initial Mass Function of young clusters from near-IR photometry. } { We consider a sample of $\\sim$300 PMS stars in the Orion Nebula Cluster (age$\\simeq$1~Myr) with well measured luminosities, temperatures and photospheric \\jb\\hb\\kb\\ photometry. This sample is used as a benchmark for testing both atmospheric and evolutionary theoretical models. } { By analyzing the photospheric colors of our sample of young stars, we find that the synthetic \\jb\\hb\\kb\\ photometry provided by theoretical spectral templates for late spectral types ($>$K6) are accurate at the level of $\\sim$0.2~mag, while colors are accurate at $\\lesssim$0.1~mag. We tabulate the intrinsic photospheric colors, appropriate for the Orion Nebula Cluster, in the range K6-M8.5. They can be conveniently used as templates for the intrinsic colors of other young (age$\\lesssim$5~Myr) stellar clusters. } { The theoretically-predicted \\jb\\hb\\kb\\ magnitudes of young late type stars do not accurately reproduce the intrinsic ones of the Orion Nebula Cluster members. An empirical correction of the atmospheric templates can fix the discrepancies between expected and observed colors. Still, other biases in the evolutionary models prevent a more robust comparison between observations and theoretical absolute magnitudes. In particular, PMS evolutionary models seem to consistently underestimate the intrinsic near-infrared flux at the very late spectral types, and this may introduce spurious features in the low-mass end of the photometrically-determined Initial Mass Function of young clusters. } ",
        "introduction": "\\label{sec:intro}  In the last two decades several generations of evolutionary models of pre-main-sequence (PMS) objects have been published with continuous improvements on the treatment of the hydrostatic structure, radiation/heat transfer and thermodynamic equilibrium. There are at least 7 different families of published PMS evolutionary calculations that have been widely circulated in machine-readable formats and that span a suitable range of stellar masses: \\citet{Swenson1994}; \\citet{Dantona97} with 1998 electronic-only update (DM98); \\citet{palla99}; \\citet{Siess00} (SDF00); \\citet{BCAH98} (BCAH98) with the subsequent extensions to the substellar regime with the COND and DUSTY models \\citep{Chabrier2000}; \\citet{Yi2003}; and \\citet{Tognelli11}. However, considerable differences still exist as the various sets of evolutionary tracks show systematic differences in the predicted masses and ages of stars in the HR diagram \\citep{Hillenbrand2004,Hillenbrand2008}.  The comparison of theoretical predictions with observations on the HR diagram requires converting observed quantities, e.g.\\ spectral type and some photometry in at least two passbands, into physical parameters like effective temperature and absolute luminosity \\citep[see e.g.][]{Hillenbrand97}. This conversion is normally done using an empirical, or semi-empirical, calibration. It is also possible to take the complementary approach of deriving observational data from theoretical predictions. For example, plotting the photometry of a young stellar cluster in a color-magnitude (CMD) or two-color diagram (2CD) may show structures (like e.g.\\ the cluster isochrone) that can be compared with the results of evolutionary and synthetic photometry calculations. Varying the model parameters one can directly visualize the effects of metallicity, effective temperature scale, surface gravity, reddening law, accretion \\citep[see e.g.][D10 hereafter]{Dario10}.  No approach is without drawbacks. On the one hand, the empirical corrections needed to derive the physical parameters of PMS stars, like e.g.\\ colors  and  bolometric corrections, may have been derived on samples of stars which may not be fully representative of the sample under examination. There is a standing tradition, for example, of using for PMS stars the intrinsic colors of Main Sequence stars \\citep[e.g.][]{Kenyon95}. The long-known difference between the spectral type vs.\\ temperature relations for dwarfs and giants represents another source of uncertainty for PMS stars, that many authors overcome (or mitigate) this uncertainty using the \\lq\\lq intermediate\\rq\\rq\\ scale of \\citet{Luhman99}, which has been found adequate for PMS stars in many studies \\citep[e.g., ][]{Dario10}. On the other hand, the systematic uncertainties of evolutionary models may combine with those of the atmospheric templates and provide erroneous results. The selection of adequate model atmospheres for PMS stars is especially critical for late type stars (M-type and later), which represent the peak of the Initial Mass Function \\citep{Bastian2010} and whose spectral energy distributions (SEDs) are dominated by broad molecular absorption features. Cool atmospheres (\\teff$\\lesssim$3000~K) host a variety of molecules and provide an environment for dust condensation, whose role in the heat and radiation transfer through the photosphere cannot be neglected \\citep{AHATS01, Allard10}. Moreover, for sub-stellar objects, convection may involve the photosphere, which therefore can no longer be treated as a system in radiative equilibrium. It is therefore crucial, particularly for cool atmospheres, to include in the synthesis codes the largest number of molecular lines in order to accurately reproduce the radiation transfer through the atmosphere.  These difficulties are aggravated by observational limitations. It is hard to obtain high quality data for a statistically significant sample of PMS stars of comparable age and distance. Even in the solar vicinity, rich and young stellar clusters are typically affected by large and inhomogeneous reddening, being still enshrouded in their parental molecular cloud \\citep{Lada03}, while the stars often show accretion excess. Older systems tend to be spatially spread and affected by membership uncertainties.  In this respect, due to its relatively low foreground extinction \\citep[$A_V\\lesssim 3$,][]{Scandariato10} and  vicinity \\citep[$d=414$~pc,][]{Menten07}, the Orion Nebula Cluster (ONC) provides a unique opportunity to analyze the intrinsic colors of its members and assess the accuracy of PMS models.  In the optical wavelength range, a recent attempt to calibrate empirically the colors of PMS stars is presented in \\citet{Dario09}. In that work, based on optical BVI photometry of the ONC, the authors show how present families of synthetic spectra fail in matching the observed colors, and present a correction based on their data. In D10, they further refine this calibration, limited to the $I$ band and 2 medium bands at $\\lambda\\sim$7700~\\AA, by explicitly calibrating colors as a function of \\teff, and by decreasing the lowest \\teff\\ limit down to $\\sim$2800~K.  The goal of this paper is to use the extensive set of spectral types and photometric data available for the ONC to test the theoretical models and to empirically determine the intrinsic (photospheric) \\jb\\hb\\kb\\ magnitudes and colors of the cluster members as functions of \\teff. The intrinsic NIR colors of PMS stars in Orion we derive will be also appropriate for the (possibly ideal) cluster isochrone, and applicable to young (age$\\lesssim$5~Myr) systems in general. In Sect.~\\ref{sec:dataset} we present the selection of our sample of stars, based on the latest spectral characterization and our recently published NIR photometry. In Sect.~\\ref{sec:analysis} we refine our list to select the subset of stars most suitable for our purposes. By means of this sample, in Sect.~\\ref{sec:atm} we test the most recent atmospheric model of \\citet{Allard10}, and in Sect.~\\ref{sec:empiric} we derive the average colors of the cluster. Finally, in Sect.~\\ref{sec:discussion} we discuss our result and we compare them to the current theoretical predictions. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have accurately analyzed a sample of $\\sim$300 stars with measured temperatures, luminosities and photospheric NIR photometry as a benchmark for current atmospheric and evolutionary models for low-mass PMS stars and brown dwarfs.  We have compared the extinction-corrected photometry to the expected photometry provided by the template spectra of \\citet{Allard10}, finding that major improvements have been done in the synthesis of theoretical spectra. Nonetheless, we obtain indications on the lack of opacity in the \\hb-band, likely due to the improper treatment of either the water vapor absorption profile or the collision induced absorption from H$_2$. We thus propose the set of empirical corrections listed in Table~\\ref{tab:atm}, to be regarded as additive terms to the synthetic NIR colors. These corrections are weakly influenced by the $\\log g$ assumed to derive synthetic colors.  We also analyzed the same sample of stars in order to derive the average isochrone of the ONC, reported in Table~\\ref{tab:colors}. The analyzed sample of stars show a magnitude spread of the order of $\\sim$3.5~mag, consistent the the $\\sim$1.5~dex luminosity spread reported in previous study of the ONC. This spread does not allow us to constrain the magnitude scale of the ONC isochrone and, by consequence, does not allow any mass determination based solely on the magnitude of stars. On the other hand, the \\jb-\\hb\\ and \\hb-\\kb\\ colors weakly depend on the photospheric luminosity, and are well constrained by our statistical analysis.  Comparing our empirical isochrone to current theoretical models, we find that there is generally good agreement with the 2~Myr old model of DUSTY and, to some extent, the 2~Myr old model of SDF00, both in magnitude (down to M5 types) and color scale. On the other hand, we find indications that the PMS isochrones of DM98 provide \\teff--$\\log L$ relations slightly flatter than our observational data.  Finally, we investigated how the theoretical models affect the photometric derivation of the IMF in the NIR domain. We find that PMS evolutionary models generally underestimate the intrinsic luminosity of VLMSs and BDs, and this may lead to artificial structures in the low-mass tail of young clusters' mass distribution.  The empirical NIR colors we have derived, in the 2MASS system, can be readily converted to other  photometric systems using the prescriptions at \\url{http://www.ipac.caltech.edu/2mass/releases/allsky/doc/sec6_4b.html}. Since low-mass stars slowly evolve in the HR diagram at very early stages, our empirical colors can conveniently be used as the intrinsic colors of young ($age\\lesssim5$~Myr) stellar clusters to derive NIR excess and extinction of individual stars.  \\begin{acknowledgement} This research has benefitted from the SpeX Prism Spectral Libraries, maintained by Adam Burgasser at \\url{http://pono.ucsd.edu/~adam/browndwarfs/spexprism}. \\end{acknowledgement}"
    },
    "1208/1208.1022_arXiv.txt": {
        "abstract": "The process of particle acceleration by left-hand, circularly polarised inertial Alfven waves (IAW) in a transversely inhomogeneous plasma is studied using 3D particle-in-cell simulation. A cylindrical tube with, transverse to the background magnetic field, inhomogeneity scale of the order of ion inertial length is considered on which IAWs with frequency $0.3 \\omega_{ci}$ are launched that are allowed to develop three wavelength. As a result time-varying parallel electric fields are generated in the density gradient regions which accelerate electrons in the parallel to magnetic field direction. Driven perpendicular electric field of IAWs also heats ions in the transverse direction. Such numerical setup is relevant for solar flaring loops and earth auroral zone. This first, 3D, fully-kinetic simulation demonstrates electron acceleration efficiency in the density inhomogeneity regions, along the magnetic field, of the order of 45\\% and ion heating, in the transverse to the magnetic field direction,  of 75\\%. The latter is a factor of two times higher than the previous 2.5D analogous study and is in accordance with solar flare particle acceleration observations. We find that the generated parallel electric field is localised in the density inhomogeneity region and rotates in the same direction and with the same angular frequency as the initially launched IAW. Our numerical simulations seem also to suggest that  the \"knee\" often found in the solar flare electron spectra can alternatively be interpreted as the Landau damping (Cerenkov resonance effect) of IAWs due to the wave-particle interactions. ",
        "introduction": "Super-thermal particles play an important role in many space plasma situations. The relevant two examples are: (i) Earths Auroral zone (AZ) that is known to host strong field-aligned currents, parallel electric field and accelerated particles. Observations essentially show two modes of particle acceleration present in AZ: a) Precipitating auroral electrons narrowly peaked at specific energy, suggestive of a static potential drop  in the AZ (e.g. Ref.\\cite{1980SSRv...27..155M}); b) More recent observations by FAST spacecraft (e.g. Ref.\\cite{2007JGRA..11205215C}) indicate existence of electrons with broad energy and narrow in pitch angle distribution that is consistent with the inertial Alfven wave (IAW)  acceleration. (ii) In solar corona, a significant fraction of the energy released during solar flares is converted into the energy of accelerated particles \\cite{2004JGRA..10910104E}. The parallel electric field that can accelerate electrons is produced when low frequency ($\\omega < \\omega_{ci}$, where $\\omega_{ci}=eB/m_i$ is the ion cyclotron frequency) dispersive Alfven wave (DAW) has a wavelength perpendicular to the background magnetic field comparable to any of the kinetic spatial scales such as: ion gyroradius at electron temperature, $\\rho_s=\\sqrt{k_B T_e/m_i}/ \\omega_{ci}$, ion thermal gyroradius, $\\rho_i=\\sqrt{k_B T_i/m_i}/ \\omega_{ci}$, \\cite{1976JGR....81.5083H} or to electron inertial length $\\lambda_e = c/ \\omega_{pe}$ \\cite{1979JGR....84.7239G}. Dispersive Alfven waves are divided into Inertial Alfven Waves or Kinetic Alfven Waves (KAW) depending on the relation between the plasma $\\beta$ and  electron/ion mass ratio $m_e/m_i$ \\cite{2000SSRv...92..423S}. When $\\beta \\ll m_e/m_i$ (i.e. when Alfven speed is much greater than electron and ion thermal speeds, $V_A \\gg v_{th,i}, v_{th,e}$) dominant mechanism for sustaining $E_\\parallel$ is the parallel electron inertia and such waves are called Inertial Alfven Waves. When $\\beta \\gg m_e/m_i$, (i.e. when  $V_A \\ll v_{th,i}, v_{th,e}$) then thermal effects become important and the dominant mechanism for supporting  $E_\\parallel$ is the parallel electron pressure gradient. Such waves are called  Kinetic Alfven Waves.  The context of this study is related to the theoretical plasma physics processes operating in the particle acceleration in solar flares. Ref.\\cite{2009A&A...508..993B}'s Introduction gives a good overview of both theoretical and observational unresolved issues. If the solar flare particle acceleration happens in the corona, in order to explain the observed X-ray flux as the electrons smash into the dense layers of the Sun, implausibly large acceleration region volumes and/or densities are required. Different ideas have been put forward: (i) Substantial re-acceleration in the chromosphere of electrons accelerated in and injected from the corona can greatly reduce the density and number of fast electrons needed to produce a X-ray flux \\cite{2009A&A...508..993B};  However, this still seems problematic because recent analysis of the evolution of a radio spectrum from a dense flare \\cite{2007ApJ...666.1256B} (Bastian et al. 2007) shows that a significant fraction of the energy in the energetic electrons can be deposited into the coronal loop, as opposed to into the chromosphere. The estimates show that the energy deposited in the corona can approach about 30\\% flare energy. (ii) The electric circuit formed by precipitating and returning electrons \\cite{2011SSRv..159..357Z,2011A&A...532A..17Z}. The effect was found from the Ampere law which handles the circuit of injected and returning electrons when simulated  with Particle-In-Cell and Fokker-Planck simulations. (iii) When wave-particle interactions in non-uniform plasma are taken into account, the evolution of the Langmuir wave spectrum towards smaller wave-numbers, leads to an effective acceleration of electrons. Thus, the time-integrated spectrum of non-thermal electrons shows an increase in super-thermal electrons, because of their acceleration by the Langmuir waves \\cite{2012A&A...539A..43K}. (iv) Solar flare triggered DAWs, as opposed to the electron beams, propagating towards the solar coronal loop foot-points, and accelerating electrons along the propagation paths \\cite{2008ApJ...675.1645F}. Key ingredient in such approach is the transverse density inhomogeneity (i.e. edges of solar coronal loops), which enable DAW to acquire non-zero perpendicular to the magnetic field wavelength, comparable to the above mentioned kinetic scales. This results in the generation of parallel electric fields, that can effectively accelerate electrons \\cite{2005A&A...435.1105T,2008PhPl...15k2902T,2011PhPl...18i2903T}.  Introduction section of Ref.\\cite{2011PhPl...18i2903T} gives an overview of the previous work on this topic in some detail. Here we mention a latest addition, Ref.\\cite{2011JGRA..11600K15M}, which studies the interaction of an isolated Alfven wave packet with a plasma density cavity. Ref.\\cite{2011PhPl...18i2903T} considered particle acceleration by DAWs in the transversely inhomogeneous plasma via full kinetic simulation particularly focusing on the effect of polarisation of the waves and different regimes (inertial and kinetic). In particular, Ref.\\cite{2011PhPl...18i2903T}  studied particle acceleration by the low frequency ($\\omega=0.3\\omega_{ci}$) DAWs, similar to considered in Ref.\\cite{2005A&A...435.1105T,2008PhPl...15k2902T}, in {\\it 2.5D geometry}, focusing on the effect of the wave polarisation, left- (ion cyclotron branch) and right- (whistler branch) circular and elliptical, in the different regimes inertial ($\\beta < m_e/m_i$) and kinetic ($\\beta > m_e/m_i$). A number of important conclusions were drawn, including (i) The fraction of accelerated electrons (along the magnetic field), in the density gradient regions is 20\\%-35\\% in 2.5D geometry. (ii) While keeping the power of injected DAWs the same in all considered numerical simulation runs, in the case of right circular, left and right elliptical polarisation DAWs with $E_y/E_z=6$ (with $x$ being the direction of the uniform background magnetic field) produce more pronounced parallel electron beams. (iii) The parallel electric field for solar flaring plasma parameters exceeds Dreicer electric field by eight orders of magnitude. (iv) Electron beam velocity has the phase velocity of the DAW. This can be understood by Landau damping of DAWs. The mechanism can readily provide electrons with few tens of keV. (v) When in 2.5D case the mass ratio was increased from $m_i/m_e=16$ to 73.44,  the fraction of accelerated electrons has increased from 20\\% to 30-35\\% (depending on DAW polarisation). This is because the velocity of the beam has shifted to lower velocity. As there are always more electrons with a smaller velocity than higher velocity in the Maxwellian distribution, one can conjecture that for the mass ratio $m_i/m_e=1836$ the fraction of accelerated electrons would be even higher than 35\\%.  In the present work we focus on the 3D effects on particle acceleration and parallel electric field generation. In particular, instead of 1D transverse, to the magnetic field, density (and temperature) inhomogeneity, we consider the 2D transverse density (and temperature) inhomogeneity in a form of a circular cross-section cylinder, in which density (and temperature) varies smoothly across the uniform magnetic field that fills entire simulation domain. Such structure mimics a solar coronal loop which is kept in total pressure balance. Section II describes the model for the numerical simulation, while the results are presented in section III. We close the paper with the conclusions in section IV. ",
        "conclusions": "The aim of this work was to explore the novelties brought about by 3D geometry effects into the problem of particle acceleration by DAWs in the solar flare and also to Earth magnetosphere auroral zone studied in an earlier 2.5D geometry work \\cite{2011PhPl...18i2903T}. It should be noted that here we have studied the case of plasma over-density, transverse to the background magnetic field, whereas references that deal with Earth auroral zone usually consider the case of plasma under-density (a cavity), as dictated by the different applications considered (solar coronal loops and Earth auroral plasma cavities). As far as the generation of parallel electric field and associated particle acceleration is concerned, there is no difference whether the transverse density gradient is positive or negative -- for the mechanism to work it has to be non-zero. This is because whistlers (and ion cyclotron waves) that propagate strictly along the magnetic field display no Landau damping, since the longitudinal component of the wave electric field is zero \\cite{zsn}, page 124. The latter becomes non-zero as the wave front turns due to the transverse density inhomogeneity. Thus we investigated a process of particle acceleration by left-hand, circularly polarised inertial Alfven wave in a transversely inhomogeneous plasma, using 3D particle-in-cell simulation. We considered a cylindrical tube that contains a transverse to the background magnetic field inhomogeneity with a scale of the order of ion inertial length. Such numerical setup is relevant for solar flaring loops and earth auroral zone. In such structure IAWs with frequency $0.3 \\omega_{ci}$ were launched  and allowed to develop three wavelength. The following key points have been established: Propagation of IAW in such a system generates time-varying parallel electric field, localised in the density gradient regions, which accelerate electrons in the parallel to magnetic field direction. Perpendicular electric field of IAW also effectively heats ions in the transverse direction. The generated parallel electric field rotates in the same direction and frequency as the \"parent\" IAW. The fully 3D kinetic simulation demonstrates electron acceleration efficiency in the density inhomogeneity regions, along the magnetic field, is of the order of 45\\% and ion heating, in the transverse to the magnetic field direction,  is about 75\\%. The latter is a factor of two times higher than the previous 2.5D analogous study and is broadly in agreement with the solar flare particle acceleration observations. Log-log plots of electron spectra seem to indicate that the \"knee\", frequently seen in the solar flare observations, can be interpreted is the Landau damping of IAWs due to the wave-particle interactions."
    },
    "1208/1208.1214_arXiv.txt": {
        "abstract": "Much uncertainty surrounds the origin of super-luminous supernovae (SNe). Motivated by the discovery of the Type Ic SN\\,2007bi, we study its  proposed association with a pair-instability SN (PISN). We compute stellar-evolution models for primordial $\\sim$\\,200\\,\\msun\\ stars, simulating the implosion/explosion due to the pair-production instability, and use them as inputs for detailed non-LTE time-dependent radiative-transfer simulations that include non-local energy deposition and non-thermal processes. We retrieve the basic morphology of PISN light curves from red-supergiant (RSG), blue-supergiant (BSG), and Wolf-Rayet (WR) star progenitors. Although we confirm that a progenitor 100\\,\\msun\\ helium core (PISN model He100) fits well the SN\\,2007bi light curve, the low ratios of its kinetic energy and \\isoni\\ mass to the ejecta mass, similar to standard core-collapse SNe, conspire to produce cool photospheres, red spectra subject to strong line blanketing, and narrow line profiles, all conflicting with SN\\,2007bi observations. He-core models of increasing  \\isoni-to-ejecta mass ratio have bluer spectra, but still too red to match SN\\,2007bi, even for model He125 -- the effect of \\isoni\\ heating is offset by the associated increase in blanketing. In contrast, the delayed injection of energy by a magnetar represents a more attractive alternative to reproduce the blue, weakly-blanketed, and broad-lined spectra of super-luminous SNe. The extra heat source is free of blanketing and is not explicitly tied to the ejecta. Experimenting with a $\\sim$\\,9\\,\\msun\\ WR-star progenitor, initially exploded to yield a $\\sim$1.6\\,B SN Ib/c ejecta but later influenced by tunable magnetar-like radiation, we produce a diversity of blue spectral morphologies  reminiscent of SN\\,2007bi, the peculiar Type Ib SN\\,2005bf, and super-luminous SN\\,2005ap-like events. ",
        "introduction": "In the last decade, a number of super-luminous supernovae (SNe) have been identified but their origin remains highly uncertain. Most of these exhibit the Type Ia/Ib/Ic SN light-curve morphology, merely ``expanded'' to form a broader and brighter peak. In this category, we find a zoo of events including, e.g., the Type IIn SN\\,2006gy \\citep{smith_etal_07a}, the peculiar type Ib SN\\,2005bf \\citep{folatelli_etal_06,maeda_etal_07} with its double-peak light curve, the type Ic SN\\,2007bi \\citep{galyam_etal_09} with its extended nebular tail, the linearly-declining SN\\,2008es \\citep{gezari_etal_09}, as well as a rather uniform group of SNe at redshifts of 0.2--1.2 with a unique, nearly-featureless,  blue continuum, and a fast fading nebular flux \\citep{quimby_etal_11}. Today, three mechanisms are proposed for these super-luminous events, representing extreme versions of SNe interacting with a circumstellar medium, \\isoni-powered SNe, and magnetar-powered SNe. For brevity, and since an interaction mechanism is believed not to be relevant to SN\\,2007bi, we restrict our discussion to the last two mechanisms.  Historically, the most natural way to explain a large SN luminosity is the production of a larger-than-average \\isoni\\ mass \\citep{colgate_mckee_69}. The essential feature is that because the half-life of \\isoni\\ is 6.075\\,d and that of its daughter nucleus \\isoco\\ is 77.23\\,d, radioactive-decay energy can reheat the ejecta once it has expanded to $\\gtrsim$10$^{14}$\\,cm and become less sensitive to $PdV$ losses. However, ``standard\" core-collapse SNe are generally ineffective \\isoni\\ producers -- the \\isoni\\ production is strongly conditioned and limited by the explosion energy and the progenitor structure. Today, highly energetic magneto-rotational explosions are expected to be the most suitable means to produce super-luminous SNe \\citep{burrows_etal_07b}. These are generally nicknamed ``hypernovae\", and are distinct from the more germane neutrino-driven core-collapse SN explosions \\citep{buras:06b}.  Pair-instability SNe (PISNe) represent an alternative for producing a large amount of \\isoni. In the exceptional instance of stars with a main-sequence mass in the range 140--260\\,\\msun, expected to form at low metallicity in the early Universe  \\citep{bromm_larson_04}, e$^{-}$e$^{+}$ pair production may lead to an explosion and give rise to a PISN \\citep{barkat_etal_67,HW02,langer_etal_07,waldman_08}. Although the explosion mechanism is robust, it is still unclear whether such massive stars can form. If they do form, their mass loss is of concern as it can considerably affect the final stellar mass and radius, and thus, the resulting explosion properties, SN radiation, and detectability \\citep{scannapieco_etal_05,kasen_etal_11}.  An alternative means to produce a bright display is by magnetar radiation \\citep{wheeler_etal_00,maeda_etal_07, kasen_bildsten_10,woosley_10}. The energy lost in the process leads to the spin down of the magnetar, which eventually quenches its power. For a dipole field, the spin-down time scale is $t_{\\rm sp}\\sim$\\,4.8$B^{-2}_{15}P^2_{10}$\\,d, where $B_{\\rm 15}$ is the magnetic-field strength in 10$^{15}$\\,G and $P_{\\rm 10}$ the rotation period $P$ in units of 10\\,ms. For suitable choices of $B_{15}$ and $P_{\\rm 10}$, this timescale can be comparable to the half-life of \\isoni/\\isoco\\ and consequently makes magnetar radiation an attractive substitute for long-lived super-luminous SNe - in combination with different ejecta masses, it also provides a natural modulation for the time to peak brightness, the luminosity at peak \\citep{kasen_bildsten_10}, as well as for the fading rate from peak.  The proposition that SN\\,2007bi is a PISN is controversial. It was discovered around the peak of the light curve ($M_R\\sim-$\\,21\\,mag), revealed a slowly fading $R$-band magnitude consistent with full \\gray\\ trapping from $\\sim$\\,5\\,\\msun\\ of \\isoni\\ (Fig.~\\ref{fig_lbol}). It exploded in an environment with a metallicity of one third solar,  which conflicts with star formation \\citep{bromm_larson_04} and evolution theory \\citep{langer_etal_07}, which expect such stars to form and explode as a Type Ic SN at much lower metallicity only. \\citet{galyam_etal_09} performed a few simulations for SN\\,2007bi, covering a range of progenitor He cores and thus explosion characteristics, and found their He100 model to be adequate. From their modeling of a nebular-phase spectrum, they infer ejecta masses compatible with the PISN scenario. Improving upon the original work of \\citet{scannapieco_etal_05}, \\citet{kasen_etal_11} studied a broad mass range of PISN progenitor models including RSG, BSG, and WR stars. Their He100 model gives a suitable match to the SN\\,2007bi light curve, as well as a rough agreement with the near-peak spectrum.  An alternative scenario, involving the collapse of a massive-star core, has been proposed by \\citet{moriya_etal_10}. In association with the extreme properties of the SN explosion they also invoke radioactive-decay energy to explain the light curve. The situation remains blurred, epitomized by the rough compatibility of both the \\isoni\\ model (PISN model He100 or extreme core-collapse SN)  and the magnetar model for explaining the SN\\,2007bi light curve \\citep{kasen_bildsten_10,kasen_etal_11}. Understanding what distinguishes these different scenarios is thus critical to identify the nature of super-luminous SNe.  In the next section we present the numerical setup for the in-depth study of PISN explosions that we have undertaken, and which will be discussed more fully elsewhere (Dessart et al., in prep). Simulation results for three different PISN progenitor models are presented in Sect.~\\ref{sect_results}. We then focus on the model  He100 that was proposed for SN\\,2007bi, and discuss its incompatibilities with the observations. Ways of alleviating these incompatibilities are discussed in Sect.~\\ref{sect_slsn}. In particular we propose two means to produce a super-luminous SN with a bluer color -- either through an increase in the \\isoni-mass to ejecta mass ratio or, alternatively, through a delayed energy injection from the compact remnant (Sect.~\\ref{sect_slsn}). ",
        "conclusions": ""
    },
    "1208/1208.4607_arXiv.txt": {
        "abstract": "\\noindent We show that, leaving aside accelerated cosmic expansion, all experimental data in high energy physics that  are commonly agreed to require physics beyond the Standard Model can be explained when completing it by three right handed neutrinos that can be searched for using {\\em current day experimental techniques}.  The model that realizes this scenario is known as Neutrino Minimal Standard Model ($\\nu$MSM). In this article we give a comprehensive summary of all known constraints in the $\\nu$MSM, along with a pedagogical introduction to the model. We present the first complete quantitative study of the parameter space of the model  where no physics beyond the $\\nu$MSM is needed to simultaneously explain neutrino oscillations, dark matter and the baryon asymmetry of the universe. This requires to track the time evolution of left and right handed neutrino abundances from hot big bang initial conditions down to temperatures below the QCD scale.  We find that the interplay of resonant amplifications, CP-violating flavor oscillations, scatterings and decays leads to a number of previously unknown constraints on the sterile neutrino properties. We furthermore re-analyze bounds from past collider experiments and big bang nucleosynthesis in the face of recent evidence for a non-zero neutrino mixing angle $\\uptheta_{13}$.  We combine all our results with existing constraints on dark matter properties from astrophysics and cosmology. Our results provide a guideline for future experimental searches for sterile neutrinos.  A summary of the constraints on sterile neutrino masses and mixings has appeared in \\cite{Canetti:2012vf}. In this article we provide all details of our calculations and give constraints on other model parameters. ",
        "introduction": "\\label{sec:intro} The Standard Model of particle physics (SM), together with the theory of general relativity (GR), allows to explain almost all phenomena observed in nature in terms of a small number of underlying principles - Poincar\\'e invariance, gauge invariance and quantum mechanics - and a handful of numbers.  In the SM these are $19$ free parameters that can be chosen as three masses for the charged leptons, six masses, three mixing angles and one CP violating phase for the quarks, three gauge couplings, two parameters in the scalar potential and the QCD vacuum angle. Three leptons, the neutrinos, remain massless in the SM and appear only with left handed chirality. GR adds another two parameters to the barcode of nature, the Planck mass and the cosmological constant.  Despite its enormous success, we know for sure that the above is not a complete theory of nature for two reasons\\footnote{We do not address theoretical issues of ``aesthetic'' nature such as fine tuning in the context of the hierarchy problem, the strong CP problem and the flavor structure. They may be interpreted as hints for new physics, but could also simply represent nature's choice of parameters.}. On one hand, it treats gravity as a classical background for the SM, which is a quantum field theory. Such description necessarily breaks down at energies near the Planck scale $M_P$ and has to  be replaced by a theory of quantum gravity. We do not address this problem here, which is of little relevance for current and near-future experiments. On the other hand, the SM fails to explain a number of experimental facts. These are neutrino oscillations, the observed baryon asymmetry of the universe (BAU), the observed dark matter (DM) and the accelerated expansion of the universe today. In addition there is a number of cosmological problems (e.g. flatness and horizon problem). These can be explained by cosmic inflation, another phase of accelerated expansion in the universe's very early history, for which the SM also cannot provide a mechanism.  To date, these are the only confirmed empirical proofs of physics beyond the SM\\footnote{We leave aside all experimental and observational anomalies that have not lead to a claim of detection of new physics, i.e. may be explained within the SM or by systematic errors.  This includes the long standing problem of the muon magnetic moment, the inconclusive results of different direct DM searches as well as various anomalies of limited statistical significance.}.  In this article we argue that, leaving aside accelerated cosmic expansion, all of them may be explained by adding three right handed (sterile) neutrinos to the SM that can be found in experiments.  The model in which this possibility can be realized is known as {\\it Neutrino minimal Standard Model} ($\\nu$MSM) \\cite{Asaka:2005an,Asaka:2005pn}. The $\\nu$MSM is an extension of the SM that aims to explain all experimental data with only minimal modifications. This in particular means that there is no modification of the gauge group, the number of fermion families remains unchanged and no new energy scale above the Fermi scale is introduced\\footnote{Because of this the technical hierarchy problem may be absent in the $\\nu$MSM because no new states with energies between the electroweak and the Planck scale are required \\cite{Bardeen:1995kv,Shaposhnikov:2007nj}.}.  The matter content is, in comparison to the SM, complemented by three right handed counterparts to the observed neutrinos. These are singlet under all gauge interactions.  Over the past years, different aspects of the $\\nu$MSM have been explored using cosmological, astrophysical and experimental constraints \\cite{Asaka:2005pn,Shaposhnikov:2008pf,Asaka:2005an,Boyarsky:2005us, Asaka:2006rw,Bezrukov:2005mx,Shaposhnikov:2006nn,Asaka:2006ek,Boyarsky:2006zi,Boyarsky:2006fg,Boyarsky:2006jm,Shaposhnikov:2006xi, Gorbunov:2007ak,Gorbunov:2007zz,Bezrukov:2007qz,Laine:2008pg, Anisimov:2008qs,Boyarsky:2009ix,Canetti:2010aw,Asaka:2010kk, Asaka:2006nq,Gorkavenko:2009vd,Asaka:2011pb,Ruchayskiy:2011aa, Gorkavenko:2012mj,Ruchayskiy:2012si,Canetti:2012vf,Drewes:2012ma}. Moreover, it was suggested that cosmic inflation \\cite{Bezrukov:2007ep,Bezrukov:2008ut,GarciaBellido:2008ab} and the current accelerated expansion \\cite{Shaposhnikov:2008xb, Shaposhnikov:2008xi,GarciaBellido:2011de,Bezrukov:2011sz} may also be accommodated in this framework by modifications in the gravitational sector, which we will not discuss here\\footnote{Inflation can be realized without modification of the gravitational interaction by adding an extra scalar to the $\\nu$MSM \\cite{Bezrukov:2009yw} (see also \\cite{Shaposhnikov:2006xi,Kusenko:2006rh,Petraki:2007gq}). This inflaton can be light enough to be detected in direct searches.}. However, though the abundances of dark and baryonic matter have been estimated individually in the framework of the $\\nu$MSM, to date it has not been verified that there is a range of right handed neutrino parameters for which they can be explained {\\it simultaneously}, in particular for experimentally accessible sterile neutrinos. In this article we present detailed results of the first complete quantitative study to identify the range of parameters that allows to simultaneously explain neutrino oscillations, the observed DM density $\\Omega_{DM}$ and the observed BAU \\cite{Canetti:2012zc}, responsible for today's remnant baryonic density $\\Omega_B$. We in the following refer to this situation, in which no physics beyond the $\\nu$MSM is required to explain these phenomena, as \\textbf{scenario I}. In this scenario DM is made of one of the right handed neutrinos, while the other two are responsible for baryogenesis and the generation of active neutrino masses. We also study systematically how the constraints relax if one allows the sterile neutrinos that compose DM to be produced by some mechanism beyond the $\\nu$MSM (\\textbf{scenario II}). Finally, we briefly comment on a {\\textbf{scenario III}, in which the $\\nu$MSM is a theory of baryogenesis and neutrino oscillations only, with no relation to DM. A more precise definition of these scenarios is given in section \\ref{sec:scenarios}. Only scenarios I and II are studied in this article, which is devoted to the $\\nu$MSM as the common origin of DM, neutrino masses and the BAU. While scenario II has previously been studied in \\cite{Canetti:2010aw}, the constraints coming from the requirement to thermally produce the observed $\\Omega_{DM}$ in scenario I are calculated for the first time in this work. We combine our results with bounds coming from big bang nucleosynthesis (BBN) and direct searches for sterile neutrinos, which we re-derived in the face of recent data from neutrino experiments (in particular $\\uptheta_{13}\\neq 0$).  Centerpiece of our analysis is the study of all lepton numbers throughout the evolution of the early universe. As will be explained below, in the $\\nu$MSM lepton asymmetries are crucial for both, baryogenesis and DM production.  We determine the time evolution of left and right handed neutrino abundances for a wide range of sterile neutrino parameters from hot big bang initial conditions at temperatures $T\\gg T_{EW}\\sim 200$ GeV down to temperatures below the QCD scale by means of effective kinetic equations. They incorporate various effects, including thermal production of sterile neutrinos from the primordial plasma, coherent oscillations, back reaction, washouts, resonant amplifications, decoherence, finite temperature corrections to the neutrino properties and the change in effective number of degrees of freedom in the SM background.  Many of these were only roughly estimated or completely neglected in previous studies. The various different time scales appearing in the problem make an analytic treatment or the use of a single CP-violating parameter impossible in most of the parameter space. Most of our results are obtained numerically. However, the parametric dependence on the experimentally relevant parameters (sterile neutrino masses and mixings) can be understood in a simple way. Furthermore, we discover a number of tuning conditions that can be understood analytically and allow to reduce the dimensionality of the parameter space.  We find that there exists a considerable fraction of the $\\nu$MSM parameter space in which the model can simultaneously explain neutrino oscillations, dark matter and the baryon asymmetry of the universe. This includes a range of masses and couplings for which the right handed neutrinos can be found in laboratory experiments \\cite{Gorbunov:2007ak}. The main results of our study, constraints on sterile neutrino masses and mixings, have previously been presented in \\cite{Canetti:2012vf}. In this article we give details of our calculation and constraints on other model parameters, which are not discussed in \\cite{Canetti:2012vf}.  The remainder of this article is organized as follows. In Section \\ref{section:numsm} we overview the $\\nu$MSM, its parametrization, and describe the universe history in its framework, including baryogenesis and dark matter production. In Section \\ref{ExistingBounds} we discuss different experimental and cosmological bounds on the properties of right-handed neutrinos in the $\\nu$MSM. In Section  \\ref{SectionKinEq} we formulate the kinetic equations which are used to follow the time evolution of sterile neutrinos and active neutrino flavors in the early universe. In Section \\ref{baryogenesissection} we present our results on baryogenesis in scenario II. In Section \\ref{DMproductionSection} we study the generation of lepton asymmetries at late times, essential for thermal dark matter production in the $\\nu$MSM. In Section \\ref{Combined} we combine the constraints of the two previous Sections and define the region of parameters where scenario I can be realized, i.e. the $\\nu$MSM explains simultaneously neutrino masses and oscillations, dark matter, and baryon asymmetry of the universe. In Section \\ref{section:concl} we present our conclusions. In a number of  appendices we give technical details on kinetic equations (\\ref{HeffSec}), on the parametrization of the $\\nu$MSM Lagrangian (\\ref{DiracConnection}), on different notations to describe lepton asymmetries (\\ref{asymmconv})  and on the decay rates of sterile neutrinos (\\ref{DecayRatesAppendix}). ",
        "conclusions": "\\label{section:concl} We tested the hypothesis that three right handed neutrinos with masses below the electroweak scale can be the common origin of the observed dark matter, the baryon asymmetry of the universe and neutrino flavor oscillations. This possibility can be realized in the $\\nu$MSM, an extension of the SM that is based on the type-I seesaw mechanism with three right handed neutrinos $N_I$. Centerpiece of our analysis is the study of sterile and active neutrino abundances in the early universe, which allows to determine the range of sterile neutrino parameters in which DM, baryogenesis and all known data from active neutrino experiments can be explained {\\it simultaneously} within the $\\nu$MSM. We combined our results with astrophysical constraints and re-analyzed bounds from past experiments in the face of recent data from neutrino oscillation experiments. We found that all these requirements can be fulfilled for a wide range of sterile neutrino masses and mixings, see figures \\ref{BA_DM_U}, \\ref{BA_SP_U} in section \\ref{Combined}. In some part of this parameter space, all three new particles may be found in experiment or observation, using upgrades to existing facilities.  This is the first complete quantitative study of the above scenario (scenario I), in which no physics beyond the $\\nu$MSM is required. We found that the $\\nu$MSM can explain all experimental data if one sterile neutrino ($N_1$), which composes the dark matter, has a mass in the keV range, while the other two ($N_{2,3}$) have quasi-degenerate masses in the GeV range. The heavier particles $N_{2,3}$ generate neutrino masses via the seesaw mechanism and create flavored lepton asymmetries from CP-violating oscillations in the early universe. These lepton asymmetries are crucial on two occasions in the early universe. On one hand they create the BAU via flavored leptogenesis. One the other hand they affect the rate of thermal DM production via the MSW effect. The second point allows to derive strong constraints on the $N_{2,3}$ properties from the requirement to explain the observed $\\Omega_{DM}$ by thermal $N_1$ production, see section \\ref{DMproductionSection}. This can be achieved by resonant production, caused by the presence of lepton asymmetries in the primordial plasma at $T\\sim 100$ MeV. The required asymmetries can be created when $N_{2,3}$ are heavier than $1-2$ GeV and the physical mass splitting between the $N_2$ and $N_3$ masses is comparable to the active neutrino mass differences. This can be achieved in a subspace of the $\\nu$MSM parameter space that is defined by fixing two of the unknown parameters (the Majorana mass splitting $\\Delta M$ and a mixing angle ${\\rm Re}\\upomega$ in the sterile sector). This choice, in which scenario I can be realized, is dubbed ``constrained $\\nu$MSM''.   We also studied systematically how the parameter constraints relax if one allows $N_1$ DM to be produced by some unspecified mechanism beyond the $\\nu$MSM (scenario II), see section \\ref{baryogenesissection}. In this case the strongest constraints come from baryogenesis and the required mass degeneracy is much weaker, $\\Delta M/M\\lesssim 10^{-3}$. We found that successful baryogenesis is possible for $N_{2,3}$ masses as low as $10$ MeV. These results are based on an extension of the analysis performed in \\cite{Canetti:2010aw} that accounts for a non-zero value of the neutrino mixing angle $\\uptheta_{13}$ and a temperature dependent Higgs expectation value. While the low mass region is severely constrained by BBN and experiments, the allowed parameter space becomes considerably bigger for masses in the GeV range. Detailed results for the allowed sterile neutrino masses and mixings are shown in figures \\ref{BA_Max_contour} - \\ref{BA_Max_mixingangle}.  If one completely drops the requirement that DM is composed of $N_1$ and considers the $\\nu$MSM as a theory of baryogenesis and neutrino oscillations only (scenario III), no degeneracy in masses is required. Note that this also implies that no degeneracy is required in scenario II if more than three right handed neutrinos are added to the SM.  For masses below $5$ GeV, the heavier sterile neutrinos can be searched for in experiments using present day technology. This makes the $\\nu$MSM one of the few truly testable theories of baryogenesis. The parameter space for the DM candidate $N_1$ is bound in all directions, see figure  \\ref{DMexclusion}, and can be tested using observations of cosmic X-rays and the large scale structure of the universe.  Since the model does not require new particle physics up to the Planck scale to be consistent with experiment, the hierarchy problem is absent in the $\\nu$MSM.     We conclude that neutrino physics can explain all confirmed detections of physics beyond the standard model except accelerated cosmic expansion.\\\\ \\newline  {\\large \\textbf{Acknowledgements}} We are grateful to Mikko Laine for providing the numerical data shown in figure \\ref{RRMplots}. We also would like to thank Oleg Ruchayskiy, Alexey Boyarsky and Artem Ivashko for sharing their expertise on experimental bounds.  This work was supported by the Swiss National Science Foundation, the Gottfried Wilhelm Leibniz program of the Deutsche Forschungsgemeinschaft, the Project of Knowledge Innovation Program   of the Chinese Academy of Sciences grant KJCX2.YW.W10 and the IMPRS-PTFS.   \\begin{appendix}     \\newpage"
    },
    "1208/1208.4577_arXiv.txt": {
        "abstract": "We introduce a new set of flow parameters to describe the time dependence of the equation of state and the speed of sound in single field cosmological models. A scale invariant power spectrum is produced if these flow parameters satisfy specific dynamical equations. We analyze the flow of these parameters and find four types of fixed points that encompass all known single field models. Moreover, near each fixed point we uncover new models where the scale invariance of the power spectrum relies on having simultaneously time varying speed of sound and equation of state. We describe several distinctive new models and discuss constraints from strong coupling and superluminality. ",
        "introduction": "Our Universe is observed to be homogeneous and isotropic at large distances from super-galactic scales ($\\sim 1$ Mpc) to the whole visible Universe ($\\sim 10^4$ Mpc). Moreover, the power spectrum of the primordial density fluctuations on such large scales is measured to be nearly scale invariant. According to the current view, these perturbations arose as quantum fluctuations in the early Universe, then left the horizon before the standard expansion phase, only to recently reenter. The evolution of the adiabatic fluctuations in the early Universe are determined by the equation of state $w$ and the speed of sound $c_s$. The question is, which types of early cosmic evolution $w(t), c_s(t)$ would give rise to the observed 10 e-folds of nearly scale invariant perturbations?  In a common class of models for generating scale invariant perturbations, the Universe is dominated by a scalar field with the canonical kinetic term and certain potential term. In this case $c_s$ equals $1$, and the evolution of the scale factor $a$, hence $H \\equiv \\dot{a}/a$ and $\\epsilon \\equiv -\\dot{H}/H^2 = \\frac{3}{2}(1+w)$, is determined by the form of the potential. Two well known such models are the inflation and the adiabatic ekpyrosis, which can both produce a scale invariant power spectrum. In the inflationary scenario \\cite{Guth:1980zm, Linde:1981mu, Albrecht:1982wi}, $\\epsilon$ stays nearly constant and close to $0$, while the scale factor $a$ grows exponentially, pushing a wide range of length scales outside the horizon. In the adiabatic ekpyrotic scenario \\cite{Khoury:2009my, Khoury:2011ii}, however, the scale factor $a$ contracts slowly while $\\epsilon$ increases rapidly, so that the horizon scale shrinks and leaves a wide range of perturbation modes outside. The scale invariant modes created in a contraction phase can be carried on to the standard expansion phase through a cosmic bounce \\cite{Tolley:2003nx, Turok:2004gb, McFadden:2005mq}.  For a constant $c_s$, the scale invariance of the power spectrum fully constrains the time dependence of the scale factor $a(t)$ \\cite{Khoury:2010gw, Baumann:2011dt}. In addition to inflation and adiabatic ekpyrosis, there is a third type of solution in which the universe goes through an ``apex'' \\cite{Khoury:2010gw}, a slow transition from expansion to contraction. Similar to adiabatic ekpyrosis, this scenario also depends on a rapidly changing $\\epsilon$ while the scale factor $a$ is nearly constant. One difference is that here the scale invariant modes are generated during the slow expansion phase before the apex, instead of during a slow contraction phase.  More generally, the speed of sound $c_s$ may vary with time, as in single field models with a noncanonical kinetic term. Several models involving a time varying $c_s$ have been considered. In the tachyacoustic expansion scenario \\cite{ArmendarizPicon:2006if, Piao:2006ja, Magueijo:2008pm, Bessada:2009ns}, $\\epsilon$ is assumed to be constant, and the scale invariant power spectrum is produced by having a suitable time dependence of the speed of sound $c_s$. Another example is considered in \\cite{Joyce:2011kh} where $c_s$ is taken to be proportional to $\\epsilon$, so that the three-point function of the curvature perturbation is also scale invariant.  With a time dependent $c_s$, the requirement of a scale invariant power spectrum alone does not fully determine the cosmic evolution. There exist infinitely many combinations of $a(t)$ and $c_s(t)$ that can all produce scale invariant perturbations. We would like to know if these different scenarios can be described and categorized into distinctive classes of cosmological models.  In this paper we present a general scheme to describe single field cosmological models by a new set of flow parameters. Any single field model can be represented by a point or a trajectory in this parameter space. The requirement of a scale invariant power spectrum constrains the dynamics of the flow parameters and determines particular classes of trajectories. We show that all existing models correspond to nearly constant values of the flow parameters. By looking for fixed points in the parameter space, we find four general types of models that encompass inflation, adiabatic ekpyrosis, apex, and tachyacoustic expansion respectively. Furthermore, we analyze the flow lines near the fixed points in different cross sections of the parameter space. By looking for various behavior of each fixed point, we obtain new cosmological models that can produce scale invariant perturbations equally well.  Our approach is very different from the conventional way of starting from a particular scalar field Lagrangian and studing its cosmological solutions. In that approach, generally speaking, only solutions in a limited region of field phase space would yield scale invariant perturbations, corresponding to a particular type of mechanism represented by some part of the flow parameter space. In our approach, however, by tracing the flow lines in the whole parameter space, we obtain all possible cosmic evolutions that can create a scale invariant power spectrum. Each and every flow line leads to scale invariant perturbations, but different trajectories may correspond to very different scalar field models. This approach effectively avoids the limitations of particular scalar field models.  In section~\\ref{sec:flow} we define the flow parameters and specify their dynamics in order to produce scale invariant perturbations. We use these parameters to study the $c_s = 1$ case and compare our results with known models. In section~\\ref{sec:stat} we generalize to a time dependent speed of sound and analyze all four fixed points of the flow parameters. In particular, we uncover new scenarios and describe the corresponding new types of cosmological models. A few examples are shown in section~\\ref{sec:exam}. Section~\\ref{sec:cons} provides further discussions on physical constraints from strong coupling and superluminality etc. We conclude with a summary in section~\\ref{sec:summ}. ",
        "conclusions": "\\label{sec:summ}  We have presented a set of flow parameters that are used to depict single field cosmological models. Scale invariant perturbations are produced in the model if these flow parameters satisfy specific dynamical equations. We have analyzed the flow of these parameters by identifying the fixed points and their properties. Besides existing models including inflation, adiabatic ekpyrosis, apex, and tachyacoustic expansion, four new scenarios have emerged from our analysis: \\begin{itemize} \\item A generalization of the adiabatic ekpyrotic model with a time varying speed of sound, described in section~\\ref{sec:adek}. In a particular model, the scale invariance of the power spectrum completely relies on the time dependence of $c_s$. \\item An extremely decelerated expansion with a decreasing speed of sound, described in section~\\ref{sec:apex}. \\item An exponentially rapid contraction, or ``deflation'', described in section~\\ref{sec:inf}. Scale invariant modes are generated when $\\epsilon$ increases and $c_s$ decreases exponentially. \\item An ``acoustic ekpyrotic'' contraction, described in section~\\ref{sec:eps-cs}. Based on the usual ekpyrotic contraction, it requires a particularly time dependent speed of sound. \\end{itemize}  These new models serve as distinctive examples in which both $\\epsilon$ and $c_s$ vary with time cooperatively to create scale invariant perturbations. The range of scale invariant modes can be estimated by considering physical constraints from sub-Planckian energy density and weak coupling of higher order perturbations. Unless subluminality is imposed, all these models can adequately account for the observed range of the scale invariant power spectrum.  The flow parameters that we used to find the new cosmological models are convenient for obtaining the scale invariant power spectrum as well as analyzing the physical constraints. We note that in \\cite{Geshnizjani:2011rm} a different set of flow parameters are defined by expanding the Hubble parameter $H$ as a function of the field $\\phi$. Those flow parameters can be used to effectively reconstruct scalar field models, but these cosmological models do not acquire scale invariant perturbations. It may be interesting to relate the two sets of flow parameters in order to construct single field models that implement the new mechanisms for generating a scale invariant power spectrum.    \\bigskip \\noindent \\textbf"
    },
    "1208/1208.3999_arXiv.txt": {
        "abstract": "The hypothesis of existence of primordial black holes with large masses ($\\geq 10^6 M_\\odot$), formed at the earliest stages of the Universe evolution, is considered in the paper. The possibility does not contradict some theories, see e.g. \\cite{BarkanaLoeb_PhysRep_2001}, and may match new observational data. In particular, this scenario of evolution could describe some peculiarities in distant galaxies and quasars. Calculations of evolution of central body mass in protogalaxies for different initial conditions are presented. It is shown that the sufficient rate of BH mass growth is not achieved in the standard scheme without complex additional assumptions. Moreover, the appearance of a primordial black hole in the epoch of primordial nucleosynthesis could significantly change the chemical composition around it. This can lead to different exotic stars with low mass and nonstandart metals enrichment. The proposed scheme is not considered as universal. On the other hand, if only tiny part of existed objects have the considered nature, it gives a unique possibility to study extremal stages of matter and fields evolution in our Universe. ",
        "introduction": "This paper considers the possibility of existence of rare high-redshift objects: galaxies and quasars, formed much earlier than the typical composition of the Universe. We also propose the mechanism of the formation of such objects.  Today the theory of formation of the structure of the Universe is well confirmed by observations. It describes the whole history of our world evolution. Observations cover only tiny part of it: they are limited now to the redshift range $0<z<10$ (see refs. \\cite{LehnertEtAl_Nature_2010,Bouwens_Nature_2011}) and the point $z\\approx 1000$ (CMB observations). Modern and future missions and experiments have the possibility to expand the region towards the moment of the Universe creation. On this way we are discovering now more and more distant galaxies with no end in sight. The theory tells us that everything around was built in evolutionary way: for some period after recombination there was no source of light in the Universe -- Dark Ages, firsts stars appeared approximately at $z\\approx 30$, and only further galaxies were constructed. So the observational trend of discovering new objects (galaxies and quasars) should be violated at some redshift. At this point we expect the discovery of something beyond standard predictions.  But what if we find a galaxy so old that it lies beyond the redshift region allowed by the standard theory (ST)? Then something in our understanding needs changing. We consider the theory to be generally correct, that is the number of such extraordinary objects is negligible. But it requires small additions to be applied. We propose primordial big black holes (PBBHs) as simple candidates to such extension. By this term we mean black holes with masses $M>M_\\odot$ formed in the early Universe (before recombination). In this case we only change the initial conditions for the problem of structure formation. The result could account for the early formation ($z_{\\rm formation}\\gg 30$) of rare objects (we propose to call them ``cosmological dinosaurs''). We state that it may be not so fantastic: recently two galaxies with unusually high black hole-to-bugle mass ratio were discovered by \\cite{BogdanFormanEtAl_1203.1641}, and a discovery of a star with very low metallicity, see \\cite{CaffauEtAl_Nature_2011}. These examples offer difficulties for the theory. They are candidates for ``dinosaurs''.  The structure of the paper is the following. In the Section \\ref{sec:st} the standard theory of the structure formation is considered along with some its bottlenecks. The Section \\ref{sec:SMBH} considers the question of supermassive black holes growth. We show that the creation of a SMBH evolutionary (and after recombination) is not a simple problem even for known today quasars. We consider their masses as a simple general criterium for the dinosaurs: too massive BH in the early Universe could not be described by the standard theory. The general aim of this article is to introduce the concept of cosmological dinosaurs and show that there could be need in some cases. Further works will consider this possibility more carefully. ",
        "conclusions": "The aim of the paper is to introduce the conception of ``cosmological dinosaurs'' -- objects appeared long before the period of intensive structure formation according to the standard model. The necessity for them could arise in future: the observations are discovering now more and more distant galaxies and quasars with no end in sight. At some moment this ``flow'' of newly opened objects could contradict the accepted theory. We consider in this case that the number of additional discoveries will be very small, so only light extension of the theory is necessary. In our work the appearance of such objects is connected with primordial big black holes (PBBHs, we propose to introduce this term as rich cosmological physics is connected with them). From our point of view these are the most natural and simple extension: they start to form the structure around themselves much earlier. The exact quantitative contribution could be calculated in more sophisticated models, than presented in this paper (e.g. with account for hydrodynamics) and is planned in future works. In the framework of our model we have shown that modern high-redshift objects are on the edge of the standard theory predictions. So the discoveries of dinosaurs or effects of their previous appearance could be made in the nearest future, introducing good probes for new physics in the early Universe.  SIG is partly supported by the project ``Development of ultrahigh sensitive receiving systems of THz wavelength range for radio astronomy and space missions'' in NSTU n.a. R.E. Alekseev, LCN."
    },
    "1208/1208.1678_arXiv.txt": {
        "abstract": "Microlensing by the stellar population of lensing galaxies provides an important opportunity to spatially resolve the accretion disc structure in strongly lensed quasars. Disc sizes estimated this way are on average larger than the predictions of the standard Shakura-Sunyaev accretion disk model. Analysing the observational data on microlensing variability allows to suggest that some fraction of lensed quasars (primarily, smaller-mass objects) are accreting in super-Eddington regime. Super-Eddington accretion leads to formation of an optically-thick envelope scattering the radiation formed in the disc. This makes the apparent disc size larger and practically independent of wavelength. In the framework of our model, it is possible to make self-consistent estimates of mass accretion rates and black hole masses for the cases when both amplification-corrected fluxes and radii are available. ",
        "introduction": "Since the work of \\citet{lyndenbell69}, disc accretion onto supermassive black holes is a commonly accepted interpretation for the activity of quasars, both radio-loud and radio-quiet (sometimes distinguished as quasi-stellar objects, QSO). Among all the active galactic nuclei, quasars are distinguished by higher luminosities (exceeding that of host galaxies) that is most likely connected to higher accretion rates.  Spectral energy distributions in optical and UV are reasonably consistent \\citep{elvis_seds} with the predictions of the standard thin accretion disc model introduced in the seminal works of \\citet{shakura72,SS73, LP74,NT73}. However, the predicted angular sizes of quasar accretion discs are too small (microarcseconds and less) to be resolved directly. For today, quasars remain essentially point-like (``quasi-stellar'') objects resolved only indirectly, in particular by microlensing effects.  As it was shown by \\citet{agolkrolik99}, microlensing by the stellar population of lensing galaxies is sensitive to the size of the emitting region. Here, we adopt the statement of \\citet{size} that the basic quantity that microlensing amplification maps and curves are sensitive to is \\emph{half-light radius}. Half-light radius $R_{1/2}$ is defined as the radius inside which half of the observed flux is emitted at a given wavelength.   Numerous studies aimed on probing the spatial properties of quasar accretion discs with help of microlensing. While most early works (see \\citet{wambsganss} for review) reported reasonable agreement between the observational data and the standard accretion disc theory, several important results are at odds with the theoretical predictions. Studying microlensing amplification statistics, \\citet{pooley07} find best-fit disc sizes more than one order of magnitude larger than the theoretical predictions based on photometrical data. Partially, this may be attributed to the mass estimates used in this study (see discussion in \\citet{paper1}). In \\citet{morgan10}, inconsistency is somewhat smaller (about a factor of 3) but still significant. Accretion discs seem too large for their apparent luminosities or too faint for their sizes in the UV/optical range ($\\sim 2000\\div 4000\\AAA$). Many papers (such as \\citet{jimenez}, \\citet{pooley07} and \\citet{morgan10}) interpret this inconsistency as an evidence for insufficiency of the standard accretion disc model, but no universal solution was proposed so far to account for the size discrepancy. There are indications for possibly higher black hole masses in some objects \\citep{morgan10,paper1}, but no changes in masses, accretion rates and efficiencies can explain the observed sizes and fluxes simultaneously.  One of the important issues in quasar microlensing studies is whether the disc radial scale $R_S$ \\citep{morgan10} dependence on wavelength is consistent with the power law $R_S \\propto \\lambda^{4/3}$ predicted by the standard accretion disc theory. Several important disc models predict power law dependences $R(\\lambda) \\propto \\lambda^\\zeta$ with different exponents. Below we will refer to $\\zeta$ as ``structure parameter''. While for \\einc, classical $\\zeta=4/3$ works fairly good \\citep{eigenbrodII,anguita08}, other objects such as \\sdss, for instance, clearly require smaller $\\zeta$. \\citet{floyd09} propose angular momentum inflow at the inner edge of the disc in \\sdss\\ \\citep{agolkrolik} as an explanation for the apparently very steep temperature law in the disc. This model implies $R_S \\propto \\lambda^{8/7}$ marginally consistent with the observational data.  In the recent work by \\citet{blackburne}, several other objects were shown to have much shallower $R(\\lambda)$ dependences, some consistent with $R_S = const$ for a broad range of comoving wavelengths, $0.1\\mu \\lesssim \\lambda \\lesssim 1\\mu$. The only object having conventional thin-disc scaling is \\mg\\ that has the highest mass among the sample of 12 objects considered by \\citet{blackburne}. All the smaller-mass ($M\\lesssim 10^9\\Msun$) black holes are characterised by $\\zeta \\sim 0\\div 0.5$.  { Microlensing effects in the X-ray range are more profound than in the optical \\citep{pooley07}. Independently of the disc structure in the optical and UV ranges, X-ray properties are more or less similar for all the objects where microlensing effects were studied in the X-ray range \\citep{einc_xrays, dai10,morgan12}. Evidently, X-ray emission comes from somewhere inside the inner $\\sim 10\\times GM/c^2$ \\citep{chen12}, and the exact mechanisms driving the formation of the X-ray continuum and lines are yet to be revealed. }  Small structure parameters originate not only for very steep temperature slopes in multi-blackbody models. For instance, $\\zeta =0$ is naturally reproduced if the brightness distribution does not depend on wavelength. This may be achieved if the accretion disc is surrounded by an envelope optically thick to Thomson scattering. Without affecting its spectral properties, scattering changes the spatial brightness distribution of the disc radiation. In general, accretion disc will increase its apparent radius and lose its intrinsic radius dependence on wavelength. A possible origin for such a scattering envelope is super-Eddington accretion that leads to formation of a Thomson-thick wind \\citep{SS73}. Since there is observational evidence that some quasars accrete in supercritical regime, especially at larger redshifts \\citep{collin}, we consider this scenario quite plausible.  In the following section \\ref{sec:sph}, we describe a simple % scattering envelope model that we use to account for the spatial properties of microlensed quasars. It will be shown that such an envelope may result from super-Eddington accretion by a supermassive black hole. In section \\ref{sec:obs} and \\ref{sec:res}, we describe the observational data and interpret them in the framework of the scattering envelope model. In section \\ref{sec:balq}, we consider the possible connection between the putative class of supercritical quasars, broad absorption line (BAL) quasars, and make conclusions in section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  Scattering envelope formed by a super-Eddington accretion disc is a plausible model for the spatial properties of the emitting regions in some lensed quasars. Large spatial sizes ($R\\sim 10^{16}\\div 10^{17}\\rm cm$) practically independent on wavelength are an expected outcome of a moderately super-Eddington ($\\mdot \\sim 10\\div 100$) mass accretion rate. Black hole masses and mass accretion rates may be determined self-consistently if both disc size and flux estimates are present.  { The small sizes of X-ray emitting regions of microlensed quasars may be explained by existence of an avoidance cone, or supercritical funnel, in the disc wind. }  Some of our super-Eddington objects are broad absorption line quasars, and at least one (\\pg) shows signatures of a mildly relativistic outflow."
    },
    "1208/1208.0953_arXiv.txt": {
        "abstract": "We calculate the expected number of type Ia supernovae (SN Ia) in the core-degenerate (CD) scenario and find it to match observations within the uncertainties of the code. In the CD scenario the super-Chandrasekhar mass white dwarf (WD) is formed at the termination of the common envelope phase from a merger of a WD companion with the hot core of a massive asymptotic giant branch (AGB) star. We use a simple population synthesis code that avoids the large uncertainties involved in estimating the final orbital separation of the common envelope evolution. Instead, we assume that systems where the core of the secondary AGB star is more massive than the WD remnant of the primary star merge at the termination of the common envelope phase. We also use a simple prescription to count systems that have strong interaction during the AGB phase, but not during the earlier red giant branch (RGB) phase. That a very simple population synthesis code that uses the basics of stellar evolution ingredients can match the observed rate of SN Ia might suggest that the CD-scenario plays a major role in forming SN Ia. ",
        "introduction": "\\label{sec:intro}  There are {{{ {four} }}}  basic theoretical scenarios for the formation of the progenitors of Type Ia supernovae (SNe Ia). The goal of these scenarios is to form a carbon-oxygen (CO) white dwarf (WD) with a mass close to and above the Chandrasekhar mass limit $M_{\\rm Ch}$. Such WDs will go through a thermonuclear detonation \\citep{HoyleFowler1960} that is observed as an SN Ia. ($i$) In the single degenerate (SD) scenario (e.g., \\citealt{Whelan1973}; \\citealt{Nomoto1982}; \\citealt{Han2004}) a WD grows in mass through accretion from a non-degenerate stellar companion. However, it seems that the WD-mass increase is very limited, e.g., \\cite{Idan2012} for a recent paper and references therein. {{{ {Ruiter et al. (2011) consider the helium-rich donor scenario (HeRS) to be in a category separate from the canonical SD scenario. In the HeRS a He-rich star, degenerate or not degenerate, transfers mass to a CO white dwarf which explodes as it approaches the Chandrasekhar mass (e.g. \\citealt{Iben1987}). Like the SD scenario, the delay time is strongly dependent on evolutionary timescales of the stars.} }}} ($ii$) In the double degenerate (DD) scenario (\\citealt{Webbink1984, Iben1984}; see \\citealt{vanKerkwijk2010} for a paper on sub-Chandrasekhar mass remnants) two WDs merge after losing energy and angular momentum through the radiation of gravitational waves \\citep{Tutukov1979}. ($iii$) In the core-degenerate (CD) scenario for the formation of SN Ia the Chandrasekhar or super-Chandrasekhar mass WD is formed at the termination of the common envelope (CE) phase or during the planetary nebula phase, from a merger of a WD companion with the hot core of a massive asymptotic giant branch (AGB) star \\citep{KashiSoker2011, IlkovSoker2012, Soker2011, Sokeretal2012}. Observations and theoretical studies cannot teach us yet whether all scenario, or only one or two of these can work (e.g., \\citealt{Livio2001, Maoz2010, Howell2011}). {{{ { ($iv$) It is likely that a sub-Chandrasekhar mass WD can explode as a SN Ia via the `double-detonation' mechanism \\citep{Woosley1994, Livne1995}.  In this model, a sub-Chandrasekhar mass WD accumulates a layer of helium-rich material on the surface, which under the right conditions can detonate, leading to a detonation near the center of the CO WD \\citep{Fink2010}. Ruiter et al. (2011) performed a population synthesis study of the double-detonation scenario considering helium-rich donors, and found a bimodal delay time distribution, the shape depending on whether non-degenerate donors ($<1 \\Gyr$) or degenerate donors ($>1 \\Gyr$) dominate. Regarding both the DD and the double-detonation scenarios, we note that sub-Chandrasekhar SN Ia explosions are supported also by the population of WD binaries in the solar neighborhood \\citep{BadenesMaoz2012}.  } }}}  In this paper we focus on the CD scenario, and to some extend we refer to the DD scenario. The merger of a WD with the core of an AGB star was studied in the past (\\citealt{Sparks1974, Livio2003, Tout2008, Mennekens2010}). \\citet{Livio2003} suggested that the merger of the WD with the AGB core leads to a SN Ia that occurs at the end of the CE phase or shortly after, and can explain the presence of hydrogen lines. This idea was further studied recently by \\cite{Sokeretal2012}. In the CD scenario the possibility of a very long time delay (up to $10^{10}$~yr) is considered as well \\citep{IlkovSoker2012}. {{{ {However, if that delay mechanism does not work, then the CD scenario can explain SN Ia only in star forming galaxies.} }}} \\cite{Mennekens2010} {{{ { and \\cite{Ruiter2008} } }}}  find in their population synthesis calculations that many systems end in a core-WD merger process, but they did not consider these to be SN Ia. Because of its rapid rotation  (e.g., \\citealt{Anand1965}; \\citealt{Ostriker1968}; \\citealt{Uenishi2003}; \\citealt{Yoon2005}), and possibly very strong central magnetic fields (e.g., \\citealt{Kundu2012, Garcia-Berro2012}), the super-Chandrasekhar WD does not explode.  There are two key ingredients in the CD scenario, in addition to the common condition that the remnant mass be $\\ga 1.4 M_\\odot$. (1) The merger should occur while the core is still large, hence hot. This limits the merger to occur within $\\sim 10^5 \\yr$ after the common envelope phase. \\citet{KashiSoker2011} showed that this condition can be met when the AGB star is massive, and some of the ejected CE gas might fall back (see also \\citealt{Soker2012}). (2) The delay between merger and explosion should be up to $\\sim 10^{10} \\yr$ if this scenario is to account for SNe Ia in old-stellar populations, in addition to SNe Ia in young stellar populations. WDs with delay times have masses in the range $\\sim 1.4-1.5 M_\\odot$, and indeed can have a long delay {{{ { (Jorge Rueda, 2012, private communication).} }}} In previous papers another ingredient was used. It asserted that the hot core is more massive than the companion cold WD. This is true in most cases, but not in all. In some rare cases this not need be the case \\citep{Sokeretal2012}, and there is a short delay to explosion.  These ingredients involve physical processes much different from those in the DD scenario, and hence make the CD a distinguished scenario, rather than a branch of the DD scenario. The merger of the core while it is still hot might prevent an early ignition of carbon \\citep{Yoon2007}, which is one of the theoretical problems of the DD scenario (e.g., \\citealt{SaioNomoto2004}). As for the delay time, in the CD scenario it is due to the spinning-down time of the WD set by magneto-dipole radiation \\citep{IlkovSoker2012}, while in the DD scenario the delay is due to the spiraling-in time due to gravitational radiation.  Stabilizing rapidly rotating super-Chandrasekhar WDs is a delicate matter (e.g., \\citealt{Boshkayev2012}, \\citealt{Yoon2004}, \\citealt{Chen2009}, and \\citealt{Hachisuetal2012a, Hachisuetal2012b} in the SD degenerate scenario). The strong magnetic fields required in the present model most likely will enforce a rigid rotation within a short time scale due the WD being a perfect conductor. The critical mass of rigidly rotating WDs is $1.48 M_\\odot$ (\\citealt{Yoon2004} and references therein). This implies that WDs more massive than $1.48 M_\\odot$ will explode in a relatively short time. The similarity of most SN Ia suggests that their progenitors indeed come from a narrow mass range. This is $\\sim 1.4-1.48 M_\\odot$ in the CD scenario. This property of the magneto-dipole radiation torque spinning-down mechanism, {{{ {that only WDs with $M_{\\rm WD} \\la 1.48 M_\\odot$ can slow down on a very long time scale (e.g., Jorge Rueda, 2012, private communication), }  }}} explains the finding that SNe Ia in older populations are less luminous (e.g., \\citealt{Howell2001, Smith2011}), and that very massive SN Ia progenitors occur in galaxies where star formation is expected, i.e., spirals and irregulars \\citep{Scalzoetal2012}.  In the present paper we estimate the number of SNe Ia from the CD scenario. Because of several large uncertainties, we perform a crude estimate of the SN Ia progenitor formation rate. In section \\ref{sec:coincident} we point to an interesting coincident of the Chandrasekhar mass with another mass in the CD and DD scenarios. The assumptions and set up of the population synthesis are given in section \\ref{sec:assumptions}. The results on the formation rate are given section \\ref{sec:results}, together with a comparison with observations, including with core-collapse supernovae (CCSNe) rate. Our summary is in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In estimating the number of SN Ia progenitors in the core-degenerate (CD) scenario we took a different approach than previous population synthesis studies in treating the common envelope (CE) phase. We assumed that most CE processes where the mass of the white dwarf remnant of the primary (initially more massive) star is less that the mass of the core,  $M_{\\rm WD1} <  M_{\\rm core}$, end in mergers. These systems have a large envelope to WD mass ratio, $\\xi \\ga 3.5$, as can be seen in the upper panel of Figure \\ref{fig:figxi1}. Such systems have a high probability for WD-core merger \\citep{Soker2012}. This approach avoids using the highly uncertain and controversial $\\alpha$-CE parameter in estimating the final orbital separation.  We required that the binary interaction is strong during the late AGB phase of the secondary star, but not during its RGB phase. We then used the ratio of maximum radii on these two phases, as presented in Figure \\ref{fig:fig1}, to find the probability for each binary system to be in the relevant orbital separation range. The probability is given in equation (\\ref{eq:xi1}), where we took the initial orbital separation of the binary system population to span 4.5 orders of magnitudes. This simple treatment avoids using the tidal interaction expression that has some uncertainties. For having a large ratio of AGB maximum radius to RGB maximum radius the initial mass of the star should be $M_{\\rm MS} > 2.3 M_\\odot$. Such stars also form WDs of masses that are about half the Chandrasekhar critical mass, such that the merger product of two such stars reaches the critical mass. This coincidence is discussed in section \\ref{sec:coincident}.  These steps, that are described in section \\ref{sec:assumptions}, leave us with two parameters. One is the mass transfer parameter from the primary to the secondary, $\\eta$, as given by equation \\ref{eq:eta1}, and the other is the binary fraction $f_B$. The likely values for these are $\\eta \\simeq 0.9$ and $f_b \\simeq 0.65$. We take a conservative approach and use $\\eta=0.8$ and $f_b=0.5-0.65$. The value of $f_B=0.5$ is used by \\cite{Nelemansetal2012} in their comparison of different studies. The criteria for a system to be counted as a SN Ia progenitor are listed in section \\ref{subsec:counting}.  The results are given in Table \\ref{tab:Table1}. The rows without a superscript $\\ast$ include all core-WD systems entering a CE phase, while the rows marked with a superscript $\\ast$ include only systems with $M_{\\rm WD1} <  M_{\\rm core}$. The latter cases are more likely to merge and are the systems considered here. These numbers should be taken very cautiously at this preliminary stage. As can be seen from Table \\ref{tab:Table1}, our very simple approach matches rate of SN Ia deduced from observations. In section \\ref{subsec:number} we also list processes we have neglected, some that will increase and other that will decrease the numbers found here. Overall, despite the large uncertainties and crude estimate, we find it very satisfactory that the simplest population synthesis we could think of yields SN Ia average rate in the CD scenario very close to observations. This match must be reexamined with more sophisticated population synthesis calculations that take mergers into account and where onset of the CE phase is reconsidered.  \\cite{Nelemansetal2012} summarize the results of population synthesis studies of the double-degenerate (DD) scenario from six groups. The results span the range of $0.2-0.6$ SN Ia per $1000 M_\\odot$ for $f_B=0.5$. These numbers are a factor of 2-10 smaller than our results. We think that two processes combine to give these lower numbers. First is the CE-$\\alpha$ prescription they use. Some systems in their calculations end at a too large orbital separation to become a SN Ia. The second one has to do with the treatment of the tidal interaction and mass transfer during the RGB and AGB phases. The tidal interaction and the formation of a CE phase when the secondary reach the upper AGB might be more significant than what is usually assumed in population synthesis studies. This is because of the unstable nature of massive AGB stars, as all secondaries are in our calculations. Such instabilities are likely to increase the number of WD that enter the CE phase with the secondary. We call for people who follow the RGB and AGB phases in the simulations to reconsider the tidal interaction when the AGB reaches the upper AGB, and alow for stronger interaction due to instabilities. Such instabilities might increase the orbital separation from where the WD remnant of the primary star enters the CE phase. We intend to study the effects of AGB instabilities in a future paper.  Future studies of the CD scenario will have to explore the exact nature and outcome of the merger process of a hot core with a cooler WD with a similar mass $0.7 M_{\\rm core} \\la M_{\\rm WD1} \\la 1.2 M_{\\rm core}$, as this is a major open question in the CD scenario. Three-dimensional simulations of this process are highly desired before the CD scenario can stand on a more solid ground.  On the observational side, if the CD scenario does indeed account for a substantial fraction of SN Ia, then rapidly rotating super-Chandrasekhar WDs should be found. Claims for massive rapidly rotating and strongly magnetized WDs have already been made, e.g. see discussion by \\citet{Malheiro2012} and \\citet{Wickramasinghe2000}. Out of the 16 magnetic WDs with well determined mass reported by \\citet{Wickramasinghe2000}, two have masses close to the Chandrasekhar limit, $M>1.3 M_\\odot$. Recent studies also show that magnetic WDs on average are more massive than non-magnetic WDs (e.g., \\citealt{Vanlandingham2005}), in particular for $M_{\\rm WD} > 1.3 M_\\odot$ \\citep{Nalezyty2004}. WD with a long time delay to explosion should be found in the mass range $\\sim 1.4-1.5 M_\\odot$ {{{ { (Jorge Rueda, 2012, private communication). } }}} More massive merger remnants are expected to explode in star forming regions, and in some cases even have massive hydrogen-rich nebula around them \\citep{Sokeretal2012}. For further discussion on the expected fraction of massive WDs see \\cite{IlkovSoker2012}. There are also theoretical studies of super-Chandrasekhar WDs stabilized by strong magnetic field (e.g., \\citealt{Kundu2012}) that should be extended to systems relevant to the CD scenario.   We thank Xiangcun Meng {{{ {and Ashley Ruiter} }}} for helpful comments. This research was supported by the Asher Fund for Space Research at the Technion, and the Israel Science foundation."
    },
    "1208/1208.2167_arXiv.txt": {
        "abstract": "{ We consider a ground array of scintillation and water Cherenkov detectors with the purpose of determining the muon content of air showers. The different response characteristics of these two types of detectors to the components of the air shower provide a way to infer their relative contributions. We use a detailed simulation to estimate the impact of parameters, such as scintillation detector size, in the determination of the size of the muon component. } ",
        "introduction": "The measurement of the mass composition of ultra-high energy cosmic rays is one of the keys that can help us elucidate their origin. Using a surface detector array to disentangle the contributions to the detector signal from the different components of an air shower is a way to accomplish this.  Up to energies around \\SI{e15}{eV}, the Galaxy is believed to be the source of cosmic rays. Several acceleration mechanisms are certainly at play but it is widely expected that the dominant one is first order Fermi acceleration at the vicinity of supernova remnant shock waves. These Galactic accelerators should theoretically become inefficient between \\SI{e15}{eV} and \\SI{e18}{eV}. The KASCADE experiment has measured the energy spectra for different mass groups in this energy range and found that there is a steepening of the individual spectra at an energy that increases with the cosmic ray mass \\cite{Apel:2008cd}. As a result, the mass composition becomes progressively heavy. It is also thought that extra-galactic sources can start to contribute to the total cosmic ray flux at energies above \\SI{4e17}{eV}. The onset of such an extra-galactic component would probably produce another change in composition.  The $X_{\\text{max}}$ measurements from the HiRes-MIA experiment have been interpreted as a change in composition, from heavy to light, starting at \\SI{4e17}{eV} and becoming proton-dominated at \\SI{1.6e18}{eV} \\cite{AbuZayyad:2000ay,Abbasi:2009nf}. Both HiRes and the Pierre Auger Observatory has measured a suppression of the flux of cosmic rays at the highest energies \\cite{Abbasi:2005ni,Abraham:2010mj} and the $X_{\\text{max}}$ measurements hint at a light or mixed composition that becomes heavier beyond \\SI{2e18}{eV} \\cite{Abraham:2010yv}.      Roughly speaking, the techniques for inferring the mass composition of cosmic rays can be split in two categories, depending on whether they exploit the sensitivity to the depth of shower maximum ($X_{\\text{max}}$) or to the ratio of the muon and electromagnetic components of the air shower \\cite{LetessierSelvon:2011dy}. Direct measurements of the fluorescence emission fall in the first category, and so do the various measurements of the Cherenkov light produced by air showers. Most ground-based detector observables depend one way or another on the number of muons in the air shower. However, the arrival time profile of shower particles has been used as an observable mostly sensitive to $X_{\\text{max}}$, in particular the so-called \\textit{rise-time}, the time it takes for the signal to rise from 10\\% to 50\\% of the integrated signal \\cite{Wahlberg:2009zz}. The measurement of the number of muons and electrons in an air shower can be done directly, for example, the way it was done with the KASCADE detector \\cite{Antoni:2003gd}.   It is important that we disentangle the contributions from the different components of the air shower, as this relates to the primary mass as well as possible systematic uncertainties arising from the use of Monte Carlo hadronic interaction generators.%  The Pierre Auger Observatory is developing a series of enhancements that aim at measuring showers in the energy range between \\SI{e17}{eV} and \\SI{e19}{eV} \\cite{amiga_2011,heat}. In particular, the objective of the AMIGA enhancement \\cite{amiga_2011} is the measurement of the muon component of the air showers using scintillators shielded by several meters of soil. In the same spirit, we are considering a \\textit{combined} surface array, consisting of two super-imposed ground arrays, a Water Cherenkov Detector (WCD) array and a scintillation detector array. The purpose of the scintillation detectors is to increase the sensitivity to the electromagnetic component of air showers.  In this article we consider the possibility of studying extensive air showers induced by cosmic rays using a combined detector consisting of water Cherenkov and scintillation detectors. In order to do this, we have developed a detailed simulation and reconstruction chain whose characteristics will be briefly described in Section \\ref{section:simulation}. We will then look at the general features that allow us to gain sensitivity to the mass composition of cosmic ray primaries at energies around \\SI{e18}{eV} and conclude, in Section \\ref{sect:e19.5_eV}, by considering the possibility of determining the contribution from the muon component of a shower from a pair of scintillation and WCD detectors. ",
        "conclusions": "We have conducted a detailed study of the sensitivity of a combined scintillator/WCD array to primary cosmic-ray mass composition. Studying the response of this array to \\SI{e18}{eV} showers, we have concluded that adding photon converters on top of the scintillators to enhance the signal from the electromagnetic component of the shower does not increase the sensitivity. We have also concluded that having many small scintillation detectors is better than having few detectors of larger size since the optimum distance is closer to the shower axis. Specifically, an array of \\SI{3.2}{m^2} detectors separated by \\SI{375}{m} is better than an array of \\SI{9.7}{m^2} separated by \\SI{750}{m}.  We have applied these ideas to the highest energies and determined that, for \\SI{e19.5}{eV} primaries, more than 30\\% of the events will have a pair of scintillator/WCD station within \\SI{800}{m} of the axis. Using this pair stations it would be possible to estimate the signal from the muon component in the WCD tanks with an uncertainty of 50\\% when the WCD station has an area of \\SI{10}{m^2} and the scintillator has an area of \\SI{1.6}{m^2}. This uncertainty can be reduced by increasing the scintillator size."
    },
    "1208/1208.2810_arXiv.txt": {
        "abstract": "We designed and constructed a special instrument to enable the determination of the stellar's spin orientation. The \\textbf{D}ifferential imag\\textbf{e} rotator for \\textbf{S}tellar \\textbf{Sp}in \\textbf{O}rien\\textbf{t}ation, \\textbf{DeSSpOt}, allows the simultaneous observations of two anti-parallel orientations of the star on the spectrum. On a  high resolution \\'echelle spectrum, the stellar rotation causes a slight line tilt visible in the spatial direction which is comparable to a rotation curve. We developed a new method, which exploits the variations in these tilts, to estimate the absolute position angle of the rotation axis. The line tilt is retrieved by a spectro-astrometric extraction of the spectrum. \\\\ In order to validate the method, we observed spectroscopic binaries with known orbital parameters. The determination of the orbital position angle is equivalent to the determination of the stellar position angle, but is easier to to detect. \\\\ \\desspot was successfully implemented on the high resolution Coud\\'e spectrograph of the Th\\\"uringer Landessternwarte Tautenburg. The observations of Capella led to the determination of the orbital position angle. Our value of $37.2^\\circ$ is in agreement with the values previously found in the literature. As such we verified that both method and instrument are valid.\\\\ ",
        "introduction": "Orientation of the rotation axis is unknown for most stars. Absolute position and inclination angles have been measured only for a handful of stars. The first method proposed for the  determination of these parameters relied on Differential Speckle Interferometry \\cite{Beckers}, and was successfully applied to Aldebaran in 1995 \\cite{Lagarde}. However the observations, made with a single aperture, lacked in signal to noise due to the short exposure, which made this technique effective only for the brightest stars. Differential Speckle interferometry evolved in the last decade into long baseline interferometry. The improvements both in instrumentation and theory have made it possible to image the shape of some stars. The measurements of the oblateness for Altair \\cite{Monnier}, Achenar \\cite{deSouza}, Vega \\cite{Peterson} or Alderamin \\cite{Zhao} permitted the determination of inclination and position angle of the spin axis. However, interferometric observations require to simultaneous use of several telescopes.  Spectro-astrometry is an alternative approach for probing structures in the milliarsecond scale, i.e. below the diffraction limit of the telescope. The method relies on the conservation of the spatial information along the slit direction, and can be used with any long slit or \\'echelle-spectrograph. By measuring the wavelength dependence of the position with respect to the photocentre along the spectrum, any asymmetries in the spectral energy distribution of the source can be followed down to the milliarcsecond scale on the resulting position spectrum. The latter outlines structures which are displaced from the continuum along this one direction. In order to constrain their shape or position, it is necessary to probe several directions thus to take several exposures with different slit orientations, or to rotate the image of the source on the slit. The method has been already successfully applied to the determination of jets\\cite{Whelan}, separation of binaries\\cite{Bailey}, and position of stellar spots\\cite{Voigt}.  In this paper we explain in \\S2 how a spectro-astrometric analysis of an \\'echelle-spectrum from a star permits the determination of the position angle. Then we describe in \\S3 the adequate instrumentation \\desspot and the first light results. Finally, in \\S4 we conclude with the realisation of a DeSSpOt-spectrograph for an use at Hamburg Observatory. ",
        "conclusions": ""
    },
    "1208/1208.1623_arXiv.txt": {
        "abstract": "{% We report optical time-resolved photometry of the CRTS transient CSS091109:035759+102943. Pronounced orbital variability with a 114\\,min period, large X-ray variability and the IR to X-ray spectral energy distribution suggest a classification as a magnetic cataclysmic binary, a likely AM Herculis star or polar. } ",
        "introduction": "Polars are magnetic cataclysmic variables (CVs) harboring a strongly magnetic white dwarf which accretes matter from a late-type main-sequence star \\citep[for a comprehensive overview see ][]{warner95}. Accretion in the strongly magnetic environment happens initially via free-falling streams which are later threadened by the magnetic field so that a quasi-radial accretion column in the vicinity of the magnetic pole(s) arises. These accretion regions are sources of hard and soft X-rays from the thermal plasma and the heated polar cap(s). Consequently, many of the $\\sim$100 now known polars were identified as optical counterparts of soft X-ray sources detected in the ROSAT all-sky survey \\citep[RASS, e.g.][]{beuschwo94}. More recently, numerous CVs have been identified spectroscopically in the SDSS \\citep[][and references therein]{szkody+11}. Despite the large collecting area of XMM-Newton and Chandra, only a few objects have been discovered with these observatories due to the low surface density of the objects.  An alternative route to discover CVs involves comprehensive photometric surveys performed with a high cadence like the Catalina Real-time Transient Survey \\citep[CRTS; ][]{drake+09}. The CRTS variable CSS091109:035759+102943 was suggested to be a magnetic cataclysmic variable in an eMail notification by the variable star network VSNET. This preliminary identification rested on significant intra-night variability and transitions from low to high optical states observed in the CRTS and the association with a cataloged X-ray source \\citep[2XMM J035758.6+102938; ][]{watson+09}.  Here we describe time-resolved follow-up photometry of CSS091109:035759+102943 (henceforth CSS091109) which, combined with archival multiwavelength data, led to the (almost) unique identification of the new variable. ",
        "conclusions": "We interpret the periodicity of 114 min detected in time-resolved photometry of CSS091109 and covering more than 800 cycles as the orbital period of a cataclysmic binary. The stable pronounced variability, the change between high and low brightness states in the optical and at X-ray wavelengths, and the overall shape of the SED are strongly suggestive of a magnetic binary, a so-called polar or AM Herculis star. Final confirmation though could be derived from time-resolved optical spectroscopy or polarimetry.  The object is apparently non-eclipsing which restricts the inclination to something less than 73\\degr. The light curve with its pronounced dip is likely shaped by absorption in an accretion stream, irradiation of the stream and the secondary star, projection effects of the stream and the accretion region on the white dwarf, and cyclotron beaming.   X-ray surveys with eROSITA will discover several $10^4$ new compact binaries \\citep{predehl+10, as12}. Massive spectroscopic surveys with e.g.~4MOST \\citep{rdj12} will reveal identification spectra and Gaia the distances for most of the CVs. Comprehensive photometric follow-up will be needed to uniquely identify the type of a cataclysmic binary and eventually to uncover the period distribution of the class, an opportunity for amateur involvement on a larger scale."
    },
    "1208/1208.4376_arXiv.txt": {
        "abstract": "The decay and annihilation cross sections of dark matter particles may depend on the value of a chameleonic scalar field that both evolves cosmologically and takes different values depending on the local matter density. This possibility introduces a separation between the physics relevant for freeze-out and that responsible for dynamics and detection in the late universe. We investigate how such dark sector interactions might be implemented in a particle physics Lagrangian and consider how current and upcoming observations and experiments bound such dark matter candidates. A specific simple model allows for an increase in the annihilation cross section by a factor of $10^6$ between freeze-out and today, while more complicated models should also allow for scattering cross sections near the astrophysical bounds. ",
        "introduction": "The particle physics properties of dark matter are important for three distinct aspects of its behavior: they determine how the initial abundance of dark matter arose, they govern how the dark matter distribution evolves and influences structure formation, and they delineate the possible ways in which dark matter may be detected. Of course, these three roles are not typically independent, since they all depend on the prescribed interactions between the dark matter particles themselves and also between dark matter and the Standard Model. These connections often provide a powerful motivation for particular dark matter candidates -- for example, the freeze-out abundance of weakly interacting massive particles points to new physics at the weak scale, which in turn leads to an attractive connection between dark matter and proposed solutions to the hierarchy problem, such as weak-scale supersymmetry.  The idea that dark matter could have interactions of astrophysically interesting magnitude has received a good amount of attention \\cite{Carlson:1992fn,deLaix:1995vi,Spergel:1999mh,Wandelt:2000ad,Firmani:2000ce,Rocha:2012,Peter:2012}, motivated in part by purported discrepancies between the standard $\\Lambda$CDM model and observations of structure on small scales (as described in \\cite{Salucci:2007tm}, for example). While most approaches of this form concentrate on giving an appreciable scattering cross-section to the dark matter, it is also interesting to consider enhanced annihilation cross sections \\cite{Kaplinghat:2000vt}.  One obstacle to simple implementations of this idea is that the required cross section for a thermal relic to obtain the right relic abundance is close to the weak scale, far too small to be relevant to dynamics in the late universe. In this paper we explore the idea that the dark matter cross section might be much larger now than it was at freeze-out, due to the evolution of a background field.  In a cosmological context, the evolution of background fields can assert a significant influence on the properties of dark matter as a function of spatial location or cosmic epoch \\cite{Casas:1991ky,Anderson:1997un,Amendola:1999er,Hoffman:2003ru,Farrar:2003uw,Bean:2007nx,Bean:2007ny,Bean:2008ac,Corasaniti:2008kx,Cohen:2008nb}. A straightforward way to achieve such effects is to invoke a light scalar field that interacts with dark matter and/or ordinary matter as well as through its own potential, and whose expectation value feeds into the dark-matter properties. A popular scenario along these lines is the ``chameleon mechanism,'' which acts to screen light, cosmologically relevant degrees of freedom to protect them from precision local tests of gravity~\\cite{Khoury:2003aq,Khoury:2003rn,Brax:2010kv,Gannouji:2010fc,Mota:2010uy}.  In this paper we investigate dark matter that interacts through a gauge symmetry with a coupling constant that depends on a chameleonlike scalar field. (The effects of chameleon vector bosons on laboratory experiments were considered in \\cite{Nelson:2008tn}.) Just as the properties of a cosmologically relevant scalar can be drastically modified in the presence of local density inhomogeneities or after evolving over cosmic time, so the interactions of dark matter may be  modified. We are able to find a model in which the late-time interaction strength is considerably higher than that at freeze-out -- although admittedly, this behavior does not seem generic.  We begin by reexamining the conventional story of dark matter freeze-out according to the Boltzmann equation, but with the additional ingredient that the dark matter properties are evolving with time. We then look at specific models featuring a Dirac dark matter particle and a U(1) gauge symmetry that is spontaneously broken, along with a chameleon scalar field. We study the cosmological evolution of this coupled system and calculate the dark matter properties, including annihilation and scattering cross sections. Finally we exhibit numerical solutions to a specific model, showing that the annihilation cross section can increase substantially during cosmic evolution. ",
        "conclusions": "In this paper we have investigated the possibility that the properties of dark matter depend crucially on the dynamics of a chameleon field -- a scalar field whose cosmological evolution depends not only on its bare potential but also on the local density of other matter (such as dark matter itself) in the Universe. We have shown that such a coupling allows the annihilation cross section (for example) of the dark matter particles to change by several orders of magnitude between freeze-out and today, while remaining consistent with all observational constraints. We have presented a general formalism to describe how this might happen and have provided a specific particle physics example, in which all relevant quantities can be calculated. While there are significant observational and theoretical constraints on models of this type, it is nevertheless possible for the cross section to evolve in such a way that there may be interesting implications for the detection of dark matter and for its dynamical effects on late-universe astrophysics.  There are, of course, other possible complications to this idea that are beyond the scope of the current paper but that provide interesting avenues for future study. One natural step is to couple our model directly to the Standard Model. One way to achieve this is to directly add the dark U(1) to the current SM gauge group~\\cite{Gondolo:2011eq}. Another possibility is to couple to the Standard Model through $U(1)$ kinetic mixing ~\\cite{Holdom:1985ag,Feng:2010zp}. This extension of our model should be able to easily accommodate the relevant particle physics constraints ~\\cite{Bjorken:2009mm,Batell:2009yf,Pospelov:2008jd,Pospelov:2008zw}, while easily allowing for decays of the dark gauge boson to Standard Model particles well before BBN. The dark matter annihilations would still be dominated by the channel $\\bar{\\psi}\\psi \\to AA$, since annihilation to Standard Model particles would be suppressed by the small coupling parameter for the U(1) mixing. However, it is a more delicate issue to decide what a natural route would be to couple the visible and dark scalar sectors, particularly with regards to coupling the chameleon to normal matter.  Finally, we did not attempt a careful analysis of the effect of late-universe inhomogeneities on the chameleon field or the dark matter properties on which it depends. In the specific models we considered, it seems as if such effects would be small, but a more careful examination is warranted."
    },
    "1208/1208.2599_arXiv.txt": {
        "abstract": "{} {We provide additional observational evidence that some Type Ia supernovae (SNe~Ia) show signatures of circumstellar interaction (CSI) with hydrogen-rich material.} {Early phase optical and near-infrared (NIR) light curves and spectroscopy of SN~2008J obtained by the {{\\em Carnegie Supernova Project}} are studied and compared to those of SNe~2002ic and 2005gj. Our NIR spectrum is the first obtained for a 2002ic-like object extending up to 2.2~$\\mu$m. A published high-resolution spectrum is used to provide insight on the circumstellar material (CSM).} {SN~2008J is found to be affected by $A_V\\sim$~1.9 mag of extinction and to closely resemble SN~2002ic. Spectral and color comparison to SNe~2002ic and 2005gj suggests $R_V$ $<$ 3.1. Spectral decomposition reveals the underlying SN emission matches a 1991T-like event and, since SN~2008J is as luminous as SN~2005gj ($V_{\\rm max}$ $=$ $-$20.3 mag), we conclude that their CSI emissions are similarly robust. The high-resolution spectrum reveals narrow emission lines produced from un-shocked gas characterized by a wind velocity of $\\sim$ 50 km~s$^{-1}$. We conclude that SN~2008J best matches an explosion of a SN~Ia that interacts with its CSM.} {} ",
        "introduction": "\\label{sec:intro} Supernova (SN)~2002ic was the first event identified as a Type Ia SN interacting with hydrogen-rich circumstellar material (CSM) \\citep{hamuy03,wood04}. Spectroscopically similar to the bright SN~1991T, the spectral energy distribution of SN~2002ic also exhibited prevalent narrow Balmer emission lines which are typically produced by SN--CSM interaction (CSI) in Type IIn core-collapse (CC) SNe. In SN~2002ic, the CSI explains not only its high peak bolometric luminosity (L$_{\\rm bol}$ $\\approx$ 3$\\times$10$^{43}$~erg~s$^{-1}$) and slow declining light curve, but also its strong, broad hydrogen, calcium, and iron features observed at late epochs \\citep{chugai04}. SN~2002ic-like events are extremely rare. To date only SN~1997cy \\citep{germany00,turatto00,hamuy03,deng04}, SN~1999E \\citep{rigon03}, SN~2005gj \\citep{aldering06,prieto07} and PTF11kx \\citep{dilday12} have been found to resemble SN~2002ic. Given the rarity of this kind of transient, and the opportunity they offer to better understand the progenitors of SNe~Ia, it is imperative to enlarge the observational sample of 2002ic-like SNe. This letter presents early-time optical and near-infrared (NIR) observations of the 2002ic-like SN~2008J obtained by the {\\em Carnegie Supernova Project} (CSP; \\citealp{hamuy06}). ",
        "conclusions": "\\label{sec:discussion} \\citet{hamuy03} and \\citet{aldering06} interpreted SNe~2002ic and 2005gj to be SNe~Ia interacting with H-rich CSM. This was based mainly on the fact that their spectra are well represented by the sum of two components: a smoothly varying continuum (produced by the CSI) and a diluted spectrum of a 1991T-like SN at comparable phase. In Fig.~\\ref{decomp} we show a similar decomposition for the 3 optical spectra of SN~2008J, where we used low-order polynomials to fit the continuum. The models match the spectra reasonably well, especially the \\ion{Si}{ii} dip at $\\sim$6200 \\AA, and the features between $\\sim$3500--5000~\\AA. However, the rounded peak at $\\sim$5600 \\AA\\ is not perfectly reproduced, although a faint feature associated with \\ion{S}{ii} is also present at the same wavelength in the spectrum of SN~1991T. As in SN~2005gj, the notch characterizing the spectrum of SN~1991T at 5300~\\AA\\ is barely detected in SN~2008J. The imperfect match is not surprising since the spectral decomposition model does not include the coupling between CSI radiation and SN envelope, which should alter the spectral features of the underlying SN. As SN~2008J appears to be as luminous as SN~2005gj (see Fig.~\\ref{absmag}), we conclude that the emission due to the CSI is also similar. The flux contribution of the underlying SN to the $V$-band maximum of both SNe is $\\sim$60\\%. Similarly to SNe~2002ic and 2005gj, if the CSM is optically thick, the interaction region should not completely cover the SN~Ia since its broad features are evident in the spectra. We note that \\citet{benetti06} argued that SN~2002ic might be better explained by a Type Ic event like SN~2004aw \\citep{taubenberger06}, which also shows the presence of \\ion{Si}{II}, rather than by the CSM-SN~Ia scenario suggested by \\citet{hamuy03}. However, in the recent case of PTF11kx, the first spectrum obtained unequivocally shows a 1991T-like SN whose spectrum was only barely affected by CSI \\citep{dilday12}. The discovery of PTF11kx therefore provides strong support to the interpretation of the other members of the SN~2002ic class as CSM--SNe Ia.  The narrow emission lines and P-Cygni profiles reveal the presence of un-shocked, radially expanding CSM photo-ionized by CSI radiation. The CSM likely originates from winds in the progenitor system, and it is characterized by a velocity $v_{w}\\approx$~50~ km~s$^{-1}$. This velocity is higher than that of red supergiant (RSG) or asymptotic giant branch (AGB) winds ($\\sim$10~km~s$^{-1}$), and much lower than that of Wolf-Rayet (WR) winds (up to 2000~km~s$^{-1}$). Post-AGB stars show wind velocities in the range of 100$-$400~km~s$^{-1}$ \\citep{kotak04}. However, we can not exclude the possibility that the precursor wind was accelerated by photoionization heating as noticed by \\citet{aldering06} for SN~2005gj. This would make the AGB or RSG wind velocities more compatible with the measured ones. Episodic nova events (like that suggested for PTF11kx) can also give rise to CSM velocities of 50$-$100~km~s$^{-1}$ \\citep{dilday12}. However, PTF11kx shows a complex CSM, with multiple wind velocities which we do not find in SN~2008J.  The broad components that we observe in bright hydrogen lines are probably related to the shock region. The absolute values of the Balmer decrement are higher for the narrow component than for the broad one, indicating that they are produced in different regions. When the CSI dominates the emission, the luminosity of the broad component of H$\\alpha$, $L$(H$\\alpha_{\\rm broad}$), is  proportional to the kinetic energy dissipated per unit time across the shock front \\citep{salamanca98}. Therefore, we roughly estimate the mass-loss rate to be $\\dot{M}\\approx3\\times10^{-3}~M_{\\odot}~$yr$^{-1}$. This is slightly lower than the mass-loss rate computed for SN~2002ic \\citep{kotak04}. Here we have adopted an efficiency factor $\\epsilon_{\\rm H{\\alpha}}=0.1$, shock and wind velocities as measured from the full-width-at-zero-intensity (FWZI) of the broad and narrow H$\\alpha$ components (6000 and 100~km~s$^{-1}$ respectively), and a measured $L$(H$\\alpha_{\\rm broad}$) $=$  1.3$\\times$10$^{41}$~erg~s$^{-1}$.  The host galaxy of SN~2008J is a bright spiral galaxy having $M_{B}=-$20.2~mag \\citep{doyle05}. Therefore SN~2008J, like PTF11kx and SN~1999E, is not located in an unusual environment as in the case of SNe~1997cy, 2002ic and 2005gj, all of which occurred in faint hosts. The significant amount of MIR-emission found by \\citet{fox11} at late times (593 days since discovery) indicates the presence of $\\sim$0.01~$M_{\\odot}$ of circumstellar dust. Late time (+380 days) dust emission was also observed in SN~2002ic \\citep{kotak04}. We conclude that SN~2008J resulted from CSI of a 1991T-like event, similar to SNe~2002ic, 2005gj and PTF11kx. The observations of these objects suggest the idea that more efficient CSI may occur leading to the misidentification of interacting SNe~Ia as SNe~IIn."
    },
    "1208/1208.0139_arXiv.txt": {
        "abstract": "A thorough analysis of multicolour CCD observations of two modulated RRab-type variables, XY~And and UZ~Vir is presented. These Blazhko stars show relatively simple light-curve modulation with the usual multiplet structures in their Fourier spectra. One additional, independent frequency with linear-combination terms of the pulsation frequency is also detected in the residual spectrum of each of the two stars. The amplitude and phase relations of the triplet components are studied in detail. Most of the epoch-independent phase differences show a slight, systematic colour dependence, however, these trends have the opposite sign in the two stars. The mean values of the global physical parameters and their changes with Blazhko phase are determined utilizing the Inverse Photometric Method (IPM). The modulation properties and the IPM results are compared for the two variables. The pulsation period of XY~And is the shortest when its pulsation amplitude is the highest, while UZ~Vir has the longest pulsation period at this phase of the modulation. Despite this opposite behaviour, the phase relations of their mean-physical-parameter variations are similar. These results are not in accord with the predictions of the Blazhko model of Stothers (2006, ApJ, 652, 643). ",
        "introduction": "In spite of the recently increased research activity in the field, we still do not have a satisfactory model for the phenomenon of the light-curve modulation of RR~Lyrae stars, the so-called Blazhko effect, that would be consistent with the wide variety of the ever increasing collection of observations. To find such a model, not only the phenomenology of the light-curve variation have to be known; the modulational changes in the atmospheric physical parameters of these stars also impose constraints on the potential Blazhko models. The brightness variations of Blazhko RR~Lyrae stars can be studied with unprecedented detail and precision through the data of {\\it CoRoT} and {\\it Kepler} space telescopes in a single photometric band, and these studies led to new explanations of the phenomenon with the detection of half-integer frequencies \\citep{bk12,kmsz}. However, to derive physical parameters, standard multicolour or spectroscopic observations are needed. Since the modulation periods of Blazhko stars span a wide range from about one week to many hundred days, detailed spectroscopic studies with sufficiently large-aperture telescopes are rather difficult to attain. Therefore, moderate-size earth-based multicolour photometric telescopes still have an important role in studying the Blazhko effect.  In the context of the Konkoly Blazhko Survey I and II \\citep{kbs1,aspc}, we have obtained multicolour light curves of bright, northern, fundamental-mode Blazhko RR~Lyrae stars. These data have already been successfully utilized to derive the variations in the atmospheric parameters of six modulated stars (see sect. 6 in \\citealt{rzl} and references therein).  According to the explanation of \\cite{stothers}, the Blazhko effect is `a direct consequence of a gradual strengthening and weakening of turbulent convection in the stellar envelope'. Though this model has been seriously criticized on theoretical grounds by \\cite{ko09,sm} and \\cite{mol}, it is also important to confront the predictions of the Stothers-model with observational results. The model predicts that the phase relation between the amplitude and period changes depends on the mean physical properties of the star and it is also in direct connection with the phase relations of the physical parameter variations during the Blazhko cycle. The present paper aims to test if there is any connection between the phase relation of the amplitude and period (phase) variations and the mean physical parameters and their variations during the Blazhko cycle as proposed by \\cite{stothers,stothers2011}. Therefore, we have selected two stars to study their Blazhko modulation in detail, XY~And ($\\alpha_{2000} = 01^{\\rm h}26^{\\rm m}42.\\!\\!^{\\rm s}43$, $\\delta_{2000} = +34{\\degr}04'06.\\!\\!^{\\prime\\prime}9$, $P=0.3987$~d, $P_{\\mathrm mod}=41.4$~d) and UZ~Vir ($\\alpha_{2000} = 13^{\\rm h}08^{\\rm m}44.\\!\\!^{\\rm s}32$, $\\delta_{2000} = +13{\\degr}24'08.\\!\\!^{\\prime\\prime}4$, $P=0.4593$~d, $P_{\\mathrm mod}=68.2$~d), which show different phase relations between their amplitude and period (phase) variations. XY~And and UZ~Vir were extensively observed in the course of the Konkoly Blazhko Survey I \\citep{kbs1}, but the data and the results of the light-curve analysis have not been published yet. ",
        "conclusions": "\\subsection{Light-curve modulation}  The Fourier spectra of the light curves of XY~And and UZ~Vir are described by multiplet frequencies ($kf_0\\pm nf_\\mathrm{m}$) and the modulation frequency ($f_\\mathrm{m}$), typical of Blazhko stars. The multiplet structures are triplets and quintuplets, with a single septuplet component in the spectrum of UZ~Vir ($6f_0-3f_\\mathrm{m}$). The phase of the rising branch of XY~And is strongly modulated, while a fix point on the rising branch of UZ~Vir (see Fig.~\\ref{fig:lc}) indicates that amplitude modulation dominates its light-cure variation. Looking at the phase variations of the maximum brightness of the two stars, similar-amplitude changes are observed as shown in Fig.~\\ref{fig:tojas}. This controversial behaviour of the phase variations of the rising branches and the maxima warns that the strength of the phase modulation at different parts of the pulsation light curve can be rather different.  The pulsation light curve of XY~And is much smoother in each Blazhko phase than that of UZ~Vir. Consequently, the light-curve solution of XY~And involves less pulsation and modulation components than that of UZ~Vir.  The modulation of UZ~Vir is of a rare type, since the amplitude of the low-frequency triplet components are significantly larger than that of the high-frequency ones. Accordingly, the light maximum goes in clockwise direction around the loop on the pulsation phase\\,--\\,brightness plane (Fig.~\\ref{fig:tojas}), while the maximum of most of the Blazhko stars goes around in the opposite direction, or the loop is degenerate.  For the first time, the amplitude and phase relations of the triplet components are also studied in detail. No common feature of the variations of the defined parameters of the two stars has been recognized. This result supports that Blazhko stars behave very individually. It is difficult, if at all possible, to find any overall, common property of the modulation. The most interesting finding of this investigation is the detection of the colour dependency of the asymmetry parameter and the phase relations. The colour dependence of these quantities is the opposite of each other in the two stars, which might be connected to their opposite-sign amplitude and phase variations during the Blazhko cycle. The colour dependence of the amplitude and phase relations of the triplets of other Blazhko stars has to be studied to decide whether such a connection really holds. As the detected colour dependence carries important physical information on the modulation, its further investigation and explanation should be an important step in the description and interpretation of the Blazhko effect.   \\subsection{Additional frequencies}  The additional frequencies in RR~Lyrae stars have been discovered only recently, as very precise and extended photometry is needed to find these low-amplitude signals.  Both XY~And and UZ~Vir show variability besides the pulsation and modulation with an additional, independent frequency. Series of four peaks appear in the Fourier spectra of both objects with separation that equals to their pulsation frequency. One of these peaks is identified as a linearly independent frequency component, while the other three members of the series are interpreted as linear combinations of this and the pulsation frequency. The frequency ratios of the possible main components of the additional frequency series ($f_0/f$) are 0.325 for XY And and 0.754, 0.430 or 0.301 for UZ Vir. The 0.301, 0.325 and 0.430 frequency ratio values cover the possible regime of the  third--sixth radial overtones. Unfortunately, the actual frequencies of the higher order radial modes on a large grid of stellar parameters have never been published. Thus we do not know whether or not the observed frequency ratios fit any higher order radial mode, indeed. However, as RR Lyrae models indicate that these higher order modes are strongly damped, therefore we consider it unlikely that radial modes were detected in these stars.  In the case of XY~And, two possible explanations of the additional-frequency series have been raised: (a) one frequency of the series might be a non-radial mode, (b) all four frequencies correspond to a secondary modulation of an extremely short, $\\sim5.01$\\,d period. For explaining the series of additional frequencies in UZ~Vir, we also raised two possibilities: (a) the frequency at 2.89\\,cd$^{-1}$ might be the first radial overtone, however, with an unusually large $f_0/f_1$ ratio of 0.754, (b) one of the two highest amplitude components of the additional frequency series corresponds to a non-radial mode. The recent detection of an additional frequency with 0.753-758 frequency ratio in the {\\it Kepler} data of RR Lyr (Moln\\'ar et al. submitted to ApJ), which has been identified with the first overtone mode, favours the first possibility in the case of UZ Vir, too.   It is of interest to examine the reported occurrences of other, additional frequencies among RR Lyrae variables. Both double mode RR Lyrae stars with extended satellite observations \\citep[AQ Leo and CoRoT ID 0101368812][]{aql,chadid} show  additional frequencies beyond the fundamental and the first overtone modes. Many Blazhko variables have one or more extra, independent frequencies in their light variation in addition to the pulsation and modulation components. Such frequencies have been reported by \\citet{v1127aql}, \\citet{poretti} and \\citet{gugg,gugg12}, see also Table~2. of \\citet{benko}. The $\\pm 12.5 f_\\mathrm{m}$ peaks around some of the pulsation components in MW~Lyr \\citep{mw1} might also be connected to an additional frequency, as suggested by \\citet{poretti}. In contrast, only two of the 19 non-modulated RRab stars observed by $Kepler$ show frequencies besides the fundamental-mode pulsation \\citep{nemec}, and no additional frequency has been detected in the only non-Blazhko $CoRoT$ RRab star analysed so far \\citep{paparo}. In the residual spectra of the 14 non-modulated RRab stars observed in the Konkoly Blazhko Survey I \\citep[see references in][]{kbs1}, no additional frequency has been found either. Although the majority of the $Kepler$ Blazhko targets still lack thorough analysis, it is clear that the occurrence rate of the appearance of additional frequencies is much higher among modulated RRab stars than in non-modulated ones. A possible explanation might be that the changing physical parameters during the Blazhko-cycle \\citep{mw2} temporarily fulfill the conditions required for the excitation of the other observed (radial or non-radial) modes, whereas the chance to satisfy the excitation conditions  is much smaller for stable-light-curve RRab stars. A similar explanation has been suggested for the appearance of a peculiar bump on the descending branch of RZ Lyr in the low-amplitude phases of its modulation \\citep{rzl}. However, the excited, additional mode of RZ Lyr is supposed to be in resonance with one of the harmonics of the main pulsation frequency, which explains why it does not appear in the Fourier spectrum as an additional frequency directly.   \\subsection{Physical-parameter changes during the modulation}  The modulational changes in the atmospheric parameters of XY~And and UZ~Vir have been derived applying the IP method. Contrary to the many differences between the Blazhko-modulation properties of the two objects, their mean physical parameters are rather similar; XY~And is only $\\sim$150\\,K hotter, $\\sim$4\\,$L_\\odot$ fainter, $\\sim$0.4\\,$R_\\odot$ smaller, $\\sim$0.01\\,$\\mathfrak{M}_\\odot$ less massive and $\\sim$0.22\\,dex more metal rich, than UZ~Vir. The changes in the mean temperature, radius and luminosity during the Blazhko cycle of the two stars do not show any significant difference either.  Both objects show strong luminosity variation during the modulation: $\\Delta L \\approx2\\,L_\\odot$ and $\\approx2.5\\,L_\\odot$ are derived for  XY~And and UZ~Vir, respectively. Similarly to all Blazhko RR Lyrae stars analysed with the IPM so far, XY~And and UZ~Vir are the brightest when the pulsation amplitude is the highest.  The radius variation is weak; the average radius is the highest around Blazhko maximum in both variables. The IPM results show only marginal temperature variation of XY~And during the Blazhko cycle, but the $T_\\mathrm{eff}$ and $(V)-(I_\\mathrm{C})$ curves together suggest that XY~And is about 10\\,--\\,20\\,K warmer at the high-amplitude phase of the modulation than when the pulsation amplitude is the lowest. The IPM results on UZ~Vir indicate only a modest but definite modulational variation of about 30\\,K in temperature. Both the radius and temperature are the highest at Blazhko maximum both in XY~And and in UZ~Vir.  The Blazhko model proposed by \\cite{stothers} postulates a changing magnetic field in the envelope of the star, which influences the parameters of convection. The model makes predictions about the phase relation between the amplitude and period modulations and between the period modulation and radius variation \\citep{stothers2011}, namely, these relations depend on the temperature of the star. For hot RRab stars, the \\cite{stothers} model predicts that the pulsation period is decreasing at high-amplitude Blazhko phase (the period and amplitude variations are in anti-phase), while for stars below the crossover temperature of about $6400$\\,K, the period is increasing at the same Blazhko phase (the changes of the two quantities are in phase). These relations are shown in the top two panels in the left columns of Fig.~\\ref{fig:ipres} for the two studied objects (the pulsation-period changes are plotted with continuous line). The two variables have similar temperatures well above the predicted crossover value, and they also have similar pulsation periods and metallicities. Still, the phase relations between their pulsation amplitude and pulsation period variation are the opposite of each other. This behaviour is also demonstrated in Fig.~\\ref{fig:tojas}, which shows that maxima of the two objects go along their modulational loops in the pulsation phase\\,--\\,brightness plane in opposite direction.  The Stothers model also expects a relation between the pulsation period and radius variation during the Blazhko cycle \\citep{stothers2011}. The model predicts anti-correlated variation between these two quantities in hot RRab stars, which turns into correlation for cooler RRab variables. However, both objects are hot, short-period RRab stars, and both have the highest radius at Blazhko maximum, while their pulsation-period variations are the opposite of each other.  These findings are not in accord with the predictions of the Stothers model \\citep{stothers,stothers2011}, which claims that two Blazhko stars with such similar properties are expected to display similar relations between their amplitude, pulsation period and radius variations. Our results show that even if the fundamental idea of the Stothers model is right, the relations between the studied properties of the Blazhko stars are more complicated than predicted by the model in its present form."
    },
    "1208/1208.1529_arXiv.txt": {
        "abstract": "We present a study on the structure of quiescent filament channels observed by Hinode/XRT and STEREO/EUVI from December  2006 to February 2009.  For 10 channels identified on the solar disk, we find that the emission on the two sides of the channel is asymmetric in both X-rays and EUV: one side has curved bright features while the other side has straight faint features. We interpret the results in terms of a magnetic flux rope model. The asymmetry in the emission is due to the variation in axial magnetic flux along the channel, which causes one polarity to turn into the flux rope, while the field lines from the other polarity are open or connected to very distant sources. For 70 channels identified by cavities at the limb, the asymmetry cannot be clearly identified. ",
        "introduction": "Filament channels are defined as regions in the chromosphere surrounding a Polarity Inversion Line (PIL) where the chromospheric fibrils are aligned with the PIL, indicating the presence of an axial magnetic field  \\citep{1971SoPh...19...59F, 1971SoPh...20..298F, 1998ASPC..150..257G, 1994ssm..work..303M}. In this paper we study filament channels on the quiet Sun. The studied channels are divided into two types: Type I channels are identified based on observed long and continuous H$\\alpha$ filaments; while Type II channels are identified according to the cavities at the limb observed by XRT.  We first select 10 Type I filament channels during November 2006 and December 2008. These channels are identified based on the H$\\alpha$ observations (mainly provided by Kanzelh\\\"{o}he Solar Observatory or Mauna Loa Solar Observatory) of long and continuous filaments. Type I channels are usually located in low-latitude active region remnants. A list of Type I channels is shown in Table 1. These channels can be divided into 5 groups, and the same channel at different solar rotations is classified into one group. Only five possible corresponding cavities are observed by XRT or TRACE, while the filament channels are at the limb. The visibility of cavity may be affected by the direction of the filament channels or some bright active regions close by.  \\begin{table}[!ht] \\caption{Type I Filament Channels.} \\smallskip \\begin{center} {\\small \\begin{tabular}{ccc|ccc} \\tableline \\noalign{\\smallskip} Channel &  & Cavity & Channel &   & Cavity\\\\ Date/Hemisphere  & Group & Date &Date/Hemisphere & Group & Date \\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} 2006-11-01/S & 1 & 10-27 & 2008-03-02/N+S & 3 & No\\\\ 2006-12-26/S & 1  & 12-19 & 2008-03-27/S & 3 & No\\\\ 2007-01-19/S & 1  &  No & 2008-04-24/S & 3& 05-02\\\\ 2007-04-30/N & 2 & No & 2008-11-29/N &4& 11-24\\\\ 2008-02-02/N+S & 3 & No & 2008-12-09/N & 5 & 12-03\\\\ \\noalign{\\smallskip} \\tableline \\end{tabular} } \\end{center} \\end{table}  \\begin{figure}[!ht] \\begin{center} \\includegraphics[scale=0.5]{SU_fig1.eps} \\end{center} \\caption{Hinode/XRT (a) and STEREO/EUVI (b) images of a filament channel on 2006 Dec 26.}\\label{fig1} \\end{figure}  Filament Channels on the quiet Sun are often observed as dark channels in X-rays and EUV, and an example is shown in Figure 1. Sometimes, sheared loops (indicated by white arrows in Figure 1a) within the filament channel are observed in X-rays, but not in EUV. Some sheared loops are stable, while others are transient structures. This result is consistent with the findings by \\citet{1998ASPC..150..171S} using Yohkoh/SXT observations.  \\begin{figure}[!ht] \\begin{center} \\includegraphics[scale=0.5]{SU_fig2.eps} \\end{center} \\caption{Hinode/XRT (a) and STEREO/EUVI (b) images of a filament channel on 2008 March 2.}\\label{fig2} \\end{figure}  We find that the structure on the two sides of the channels is asymmetric in both X-rays and EUV. An example of a filament channel observed on 2008 March 2  is shown in Figure 2. This figure shows that the eastern side (black arrow) has curved bright features, while the western side (white arrows) has straight faint features. This asymmetric structure is found for all of the Type I filament channels.   \\begin{figure}[!ht] \\begin{center} \\includegraphics[scale=0.5]{SU_fig3.eps} \\end{center} \\caption{Hinode/XRT images of two Type II filament channels. The white contours refer to the corresponding H$\\alpha$ filament.}\\label{fig3} \\end{figure}  We then identify 70 type II filament channels based on cavities in XRT synoptic observations, which are taken twice a day normally. First we look for cavities from these synoptic images.  Then the filament channels are identified as the disk counterpart from the images taken 7 days earlier and later, for the cavities on the west and east limb, respectively.  The H$\\alpha$ filaments in these channels are shown as dotted or short segments.  The emission on the two sides of the channels is weaker than that in Type I channels. Based on the structure on the two sides of the channel, the 16 well observed type II channels can be divided into two types: curved on the disk side and straight or unclear on the polar side (Figure 3a); straight or unclear on the disk side and straight on the polar side (Figure 3b). Only one out of the 16 well observed channels appears to show asymmetric structure on the two sides of the channel (see Figure 3a). ",
        "conclusions": ""
    },
    "1208/1208.4283_arXiv.txt": {
        "abstract": "A compact formula for the stress tensor inside a self-gravitating, triaxial ellipsoid in an arbitrary rotation state is given. It contains no singularity in the incompressible medium limit. The stress tensor and the quality factor model are used to derive a solution for the energy dissipation resulting in the damping (short axis mode) or excitation (long axis) of wobbling. In the limit of an ellipsoid of revolution, we compare our solution with earlier ones and show that, with appropriate corrections, the differences in damping times estimates are much smaller than it has been claimed.\\\\ This version implements corrections of misprints found in the MNRAS published text. ",
        "introduction": "Most asteroids rotate in the principal, shortest axis mode: their spin axes practically coincide with the directions of the maximum moment of inertia. Only 45 out of almost 5500 entries of the LCDB light curve database \\citep[March 2012 version]{Warner:09} refer to objects that are possible non-principal axis (NPA) rotators, also known as `tumblers' or wobbling objects. With one exception of 253 Mathilde, tumblers are rather small, with estimated diameters below 20~km, but even in this size range they belong to a minority among about 2000 objects of this size with known rotation periods.  Attitude dynamics of asteroids is shaped mainly by gravitational torques (exerted either systematically by the Sun and giant planets, or sporadically during encounters with other bodies), collisions, optical and thermal radiation recoil torques, i.e. the Yarkovsky-O'Keefe-Radzievskii-Paddack (YORP) effect, and -- last but not least -- by energy dissipation due to inelastic deformations. As far as NPA rotation is concerned, collisions and close approaches trigger tumbling \\citep{SOWAH:2000,PBW:02}. Small fragments created from collisions of larger objects are also expected to start their lives in a NPA rotation state. The YORP effect also excites wobbling \\citep{Rub:2000,VBNB:07,BRV:2011}, whereas -- save for possible resonances -- distant bodies gravitation torques are neutral in this respect. Thus, even accounting for observational selection effects mentioned by \\citet{Pravec:05}, the dissipative damping seems to override other effects in most of cases.  The mechanism of wobble  damping was first identified by \\citet{Prend:58}. In NPA rotation, the centrifugal acceleration oscillates periodically, deforming each body fragment. The deformation is not perfectly elastic, so some fraction of fluctuating strain-stress energy is dissipated during each precession period and converted into heat. Draining the elastic energy affects the kinetic energy of rotation which also decreases. Thus the rotation axis is driven towards the minimum energy state -- rotation around the principal axis of maximum inertia. The angular momentum, however, is not affected by the energy dissipation, as far as we ignore thermal radiation and consider the body as an isolated system. Prendergast provided a general form of energy dissipation rate equation for an oblate spheroid\\footnote{In this paper we use the word `spheroid' for an arbitrary ellipsoid of revolution.} based upon the solution of 3D elasticity equations and the assumption that a constant fraction of the oscillating part of elastic energy is dissipated at each precession period. The latter assumption defines the now commonly adopted `$Q$-model'.  \\citet{BuSaf:73} built upon the general idea of Prendergast using combination of a spheroidal shape for rotation and a bent slender beam approximation for elastic energy. Their simple estimate of spin axis alignment time is still in use -- sometimes in the version provided by \\citet{Harris:94}. However, some scepticism towards it has been brought by observations of asteroids that do rotate around the principal axis in spite of having Burns-Safronov damping time estimate longer than the age of the Solar System. Meanwhile, the problem migrated to geophysics (e.g. Chandler wobble damping), rotation dynamics of comets and interstellar dust grains physics. The last branch, stemming from \\citet{Purc:1979}, was finally brought back to the dynamics of comets and asteroids with the sequence of papers by Efroimsky and Lazarian \\citep{LazEf:1999,Efroimsky:2000,EfLaz:2000,Efroimsky:2001,Efroimsky:2002}. Their main point of novelty is the attempt to discuss a triaxial object, represented by a rectangular prism (brick), by solving the complete, quasi-static stress tensor equation \\citep{Efroimsky:2000}. Later on, \\citet{MoMo:03} issued the damping model for a spheroid using the same starting point as \\citet{Prend:58}, i.e. solving equations for displacements. The work of \\citet{SBH:05} not only provides the solution for a spheroid with two different ways of estimating the peak elastic energy required for the $Q$-model, but it also offers a long discussion of shortcomings and problems related with earlier papers mentioned in this paragraph.  Trying to combine the YORP effect with a damping mechanism, we first intended to use the spheroid based  model of \\citet{SBH:05} for arbitrary shape asteroids. This approach, mentioned in \\citet{VBNB:07}, became less appealing after a closer inspection, because of substantial difference in the dynamics of bodies with and without axial symmetry. On the other hand, the solution of \\citet{Efroimsky:2000}, albeit referring to  a triaxial shape, exhibits a number of drawbacks: \\begin{enumerate} \\item As a consequence of using a non-smooth, brick-shaped object, the solution of stress equations is inexact, with unknown error bounds. \\item Compatibility conditions are not fulfilled, i.e. there is no displacements field that might produce the strain tensor found by Efroimsky \\citep{SBH:05}. \\item Rotation dynamics is treated by approximate formulae valid only in the neighborhood of principal axis. \\end{enumerate} Later on, \\citet{Efroimsky:2001} suggested Fourier series involving the Jacobi nome as a remedy for the last item, but none of subsequent works has implemented this guideline. In these circumstances, we have decided to resume the problem at the point where Efroimsky has abandoned it, not only using Fourier series to resolve the last problem, but also applying the triaxial ellipsoid shape which resolves the first two objections as well.  From this point of view, the present work combines a stress solution in the style of \\citet{SBH:05} with energy dissipation treatment in the spirit of \\citet{Efroimsky:2001}.  In Section~\\ref{Sec:2} we first formulate the problem of determining the stress tensor and enumerate the assumptions, hoping to help a reader less familiar with elasticity problems. Basic facts are recalled according the textbooks of \\citet{LaLi:7}, \\citet{Saad:05}, and \\citet{Wilma}. Two independent methods (displacements approach and stress approach) are used to derive and cross-check the final expressions of the stress tensor.  The $Q$-model of energy dissipation is introduced in Section~\\ref{Sec:3} and applied in Section~\\ref{Sec:4} to derive an energy dissipation rate formula. Section~\\ref{Sec:5} presents wobble damping time equations based upon the results of Section~\\ref{Sec:4} and some exemplary results. In Section~\\ref{Sec:6} we present the reduction to a specific case of a spheroid where a comparison of our solution with those reported in earlier works is possible. We use this opportunity to resolve controversies concerning drastically different energy dissipation rates in various models. ",
        "conclusions": "Using the ensemble of standard assumptions, we have derived the stress tensor inside a freely rotating and self-gravitating ellipsoid. Writing about known solutions to this problem, \\citet{WaSch:2002} put meaningful quotation marks around the word \\textit{available}. Our present solution, given in the Appendix~\\ref{ap:A} has a form which is probably compact enough  to suppress the marks, especially for the principal axis rotation. Interestingly, the presented form of $\\mtr{T}$ does not involve singularity at the incompressible limit $\\nu=\\frac{1}{2}$, where \\citet{WaSch:2002} had to use $\\nu = 0.499$.  The stress tensor has served us as the basis for the energy dissipation model built along the lines that \\citet{Efroimsky:2002} proposed, but left unaccomplished. However,  the use of an ellipsoid instead of an Efroimsky-Lazarian rectangular prism permitted to avoid all objections related with partially satisfied boundary conditions and/or missing compatibility conditions. We have also found no reasons to impose superficial conditions of a stress-free boundary like \\citet{MoMo:03}.  The solution hinges upon the use of $P_k(\\mathrm{q})$ series the Jacobi nome $\\mathrm{q}$. Their convergence is very good and even taking $p \\leqslant 4$ in equations (\\ref{P:1}-\\ref{P:4}), guarantees at least three significant digits in the area under $\\Psi_s(\\theta_s)$ for $0 < \\theta_s < \\pi/2$. Of course, the series are not legitimate exactly at $\\mathrm{q}_s=1$ (i.e $\\theta_s = 90^\\circ$). An in-depth discussion of this limit was given by \\citet{Efroimsky:2001}. On the other hand,  this state is not to be considered seriously, since any additional torque will trigger the emergence of a chaotic zone in the vicinity of separatrices.  In contrast to the results of \\citet{MoMo:03}, we find that the role of compressibility in the energy dissipation process is marginal in the range of Poisson's ratio $0 \\leqslant \\nu \\leqslant \\frac{1}{4}$. Apparently, the stronger dependence obtained by \\citet{MoMo:03} for a spheroid resulted from too strong boundary conditions.  Investigating the spheroid as a particular case of our model, we have succeeded to resolve a major part of controversies concerning short damping times of \\citet{EfLaz:2000} and very long ones according to \\citet{SBH:05}. In our opinion, the excess of energy dissipation rate over mainstream models is mostly due to a higher volume of the body shape assumed by \\citet{EfLaz:2000}. The shape factors reported by \\citet{SBH:05} are overestimated mostly by an incompatible quality factor definition, a spurious factor $\\pi$ in their computations, and a missing second mode multiplier. We find the objections against boundary conditions used or compatibility conditions violation, raised by \\citet{MoMo:03} or \\citet{SBH:05}, formally justified, yet we note that they affect the accuracy of the damping/excitation times by at most $50\\%$."
    },
    "1208/1208.0069.txt": {
        "abstract": "We present a detailed study of the X-ray, optical and radio emission from the jet, lobes and core of the quasar PKS~2101--490 as revealed by new \\emph{Chandra}, HST and ATCA images. We extract the radio to X-ray spectral energy distributions from seven regions of the 13$\\arcsec$ jet, and model the jet X-ray emission in terms of Doppler beamed inverse Compton scattering of the cosmic microwave background (IC/CMB) for a jet in a state of equipartition between particle and magnetic field energy densities. This model implies that the jet remains highly relativistic hundreds of kpc from the nucleus, with a bulk Lorentz factor $\\Gamma \\sim 6$ and magnetic field of order 30~$\\mu$G. We detect an apparent radiative cooling break in the synchrotron spectrum of one of the jet knots, and are able to interpret this in terms of a standard one-zone continuous injection model, based on jet parameters derived from the IC/CMB model. However, we note apparent substructure in the bright optical knot in one of the HST bands. We confront the IC/CMB model with independent estimates of the jet power, and find that the IC/CMB model jet power is consistent with the independent estimates, provided that the minimum electron Lorentz factor $\\gamma_{\\rm min} \\gtrsim 50$, and the knots are significantly longer than the jet width, as implied by de-projection of the observed knot lengths. ",
        "introduction": "\\label{sec:intro}  The first \\emph{Chandra} observations of the quasar PKS~0637--752 revealed a bright X-ray jet extending 12$\\arcsec$ from the quasar core ($>$ 500 kpc de-projected, assuming a jet viewing angle $\\theta < 9^{\\circ}$, as evidence by the observed proper motions of the pc-scale jet using a modern cosmology \\citep{lovell00}), associated with the previously known radio jet, but with an unexpectedly high X-ray to radio flux density ratio \\citep{schwartz00, chartas00}. Since then, tens of quasar jets have been found to possess X-ray jets with similarly high X-ray to radio flux density ratios \\citep[e.g.][]{harris02, sambruna02, sambruna04, marshall05, marshall11, kataoka05, massaro11}. The strong X-ray emission from kpc-scale quasar jets such as that of PKS~0637--752 is hard to explain in terms of standard emission mechanisms such as thermal bremsstrahlung or synchrotron self Compton emission \\citep{chartas00, schwartz00}. A popular explanation for the strong X-ray emission is the beamed, equipartition IC/CMB model proposed for PKS~0637--752 by \\citet{tavecchio00} and \\citet{celotti01}, in which the flow velocity is assumed to be highly relativistic and directed close to the line of sight. A relativistic jet velocity increases the energy density of the cosmic microwave background (CMB) in the rest frame of the jet plasma, thereby increasing the X-ray emissivity produced via inverse Compton scattering of CMB photons. The small jet viewing angle implies that the emission is Doppler beamed towards the observer. The appeal of this model is largely due to its simplicity and consistency with equipartition between magnetic and particle energy densities in the emitting plasma. From here on, we refer to the beamed, equipartition IC/CMB model as simply the IC/CMB model.  A number of uncertainties and potential problems for the IC/CMB model have been identified: (1) There is no conclusive theoretical or empirical justification for the assumption of equipartition in jet plasma, although, for a given jet speed, the equipartition condition minimises the plasma energy density and jet power.  (2) The IC/CMB model requires jet Lorentz factors of order $\\Gamma \\sim 5 - 25$ on scales of hundreds of kpc from the core. Such large jet Lorentz factors are inconsistent with studies of the radio emission from large-scale jets and counter-jets \\citep[$\\Gamma \\lesssim 1.5$][]{wardle97, mullin09}. A suggested solution to this problem invokes velocity structure across the jet --- the so-called ``spine-sheath\" model, in which the radio emission from jets in FRII radio galaxies is dominated by a slower moving sheath, whilst the emission from quasar jets with small viewing angles is said to be dominated by the Doppler boosted radiation from a fast moving spine \\citep[see e.g.][]{hardcastle06, mullin09}. (3) It has been argued that, due to the long cooling timescale for the $\\gamma \\sim 100$ electrons responsible for the IC/CMB X-ray emission, radiative and adiabatic losses alone cannot account for the rapid drop in X-ray surface brightness outside the knots \\citep{tavecchio03, stawarz04a, siemiginowska07}. (4) The IC/CMB model makes strong, testable predictions about the redshift dependence of kpc-scale X-ray jets. Specifically, the model predicts that the X-ray surface brightness should be  redshift independent, because the CMB energy density increases as (1+z)$^4$ which balances the usual (1+z)$^{-4}$ decrease of surface brightness with redshift. Therefore, the ratio of X-ray to radio surface brightness should be a strong function of redshift \\citep{schwartz02, marshall11}. So far, the search for the predicted redshift dependence has been unsuccessful \\citep{marshall11, kataoka05}. (5) It has been argued \\citep[e.g.][]{atoyan04} that the very large jet powers derived from the IC/CMB model ($\\gtrsim 10^{48}$ ergs/s, e.g.\\, \\citet{tavecchio00}) are prohibitively large. Such high jet power is disfavoured because $10^{48}$~ergs/s is equal to the Eddington luminosity of a $10^{10}$ M$_\\odot$ black hole, and such high jet power is an outlier when compared to samples of FRII radio galaxies such as the \\citet{rawlings91} sample, for which the largest estimated jet power is of order $10^{47}$~ergs/s. (6) A number of sources show a decreasing X-ray to radio flux density ratio along the jet, which, if the IC/CMB model is correct, implies deceleration must be taking place on hundreds of kpc-scales \\citep{georganopoulos04}. However, it is not clear how the gradual deceleration can occur, and there is no evidence for deceleration in lobe-dominated radio galaxies, which may be expected in such a model \\citep{hardcastle06}.   Despite the numerous concerns surrounding the IC/CMB model, none of the issues listed above is currently seen as definitively refuting the model, and it continues to receive attention in the literature as the likely candidate for the quasar jet X-ray emission mechanism. In this paper, we critically assess the application of the IC/CMB model to jet X-ray emission of PKS~2101--490.   PKS 2101-490 was first reported as a bright flat spectrum radio source by \\citet{ekers69}. \\citet{marshall05} reported a redshift of z $\\approx$ 1.04 for this source, based on an unpublished Magellan spectrum (see also the discussion in \\S \\ref{sec:core}). The spectroscopic redshift determination is robust, with uncertainty of approximately $\\pm$0.003. Further details of the spectroscopic observations and data analysis will be presented in an upcoming paper (Gelbord \\& Marshall, in prep.).  Studies at the Australia Telescope Compact Array (ATCA) revealed significant radio emission on arcsecond scales \\citep{lovell97}. For this reason, PKS 2101--490 was included in the \\emph{Chandra} snapshot survey of flat spectrum radio quasars with arcsecond scale radio jets \\citep{marshall05}. \\citet{marshall05} presented an 8.6~GHz ATCA image along with a 5 kilosecond snapshot \\emph{Chandra} image (\\dataset [ADS/Sa.CXO#obs/03126] {Chandra ObsID 3126}) that revealed significant X-ray emission associated with the 13$\\arcsec$ eastern radio jet. Based on the results of the snapshot survey, and its morphological similarity to PKS~0637--752, PKS~2101--490 was selected along with a number of other sources: PKS~1421-490 \\citep{godfrey09, gelbord05}; PKS~1055+201 \\citep{schwartz06b}; PKS~0208-512 \\citep{perlman11}; PKS~1202-262 \\citep{perlman11}; PKS~0920-397 \\citep{schwartz10}; and PKS~1030-357, as a target for deeper follow-up observations with \\emph{Chandra}, HST and ATCA. Here we present a detailed physical analysis of the jet, hotspot and lobes of this source based on new ATCA, \\emph{Chandra} and HST follow-up images.    \\begin{deluxetable}{cccc} \\tabletypesize{\\scriptsize} \\tablecaption{Observation Information \\label{table:2101_obs_info}} \\tablewidth{0pt} \\tablehead{ \\colhead{Instrument} & \\colhead{Band} & \\colhead{Mode} & \\colhead{Date} } \\startdata ATCA & 4.8~GHz & 1.5A/6A & May 25/Sep 2 2000 \\\\ ATCA & 8.64~GHz & 1.5A/6A & May 25/Sep 2 2000  \\\\ ATCA & 17.73~GHz & 6C & May 10 2004 \\\\ ATCA & 20.16~GHz & 6C & May 10 2004  \\\\ HST & F814W & ACS/WFC & Mar 8 2005  \\\\ HST & F475W & ACS/WFC & Mar 8 2005  \\\\ \\emph{Chandra} & $0.5-7$~keV & ACIS-S3 & Dec 17 2004 \\enddata \\end{deluxetable}    In \\S \\ref{sec:observations} we describe the observations and data reduction. In \\S \\ref{sec:results} we discuss the method and results of modelling the spectral energy distributions of spatially separated jet knots in terms of the IC/CMB model. In \\S \\ref{sec:jet_power} we compare independent jet power estimates with that obtained from the IC/CMB model of jet X-ray emission. In \\S \\ref{sec:2101_K6_broken_powerlaw} we discuss the optical emission detected from one of the jet knots, and present an interpretation of this in terms of a broken power law synchrotron spectrum. In \\S \\ref{sec:core} we discuss the X-ray spectrum of the quasar core. In \\S \\ref{sec:lobes} we discuss the radio and X-ray emission from the lobes, in particular, we model the lobe SEDs in terms of inverse Compton scattering of the CMB. In \\S \\ref{sec:conclusions} we present the conclusions and final remarks. ",
        "conclusions": "\\label{sec:conclusions}  We have presented an analysis of new \\emph{Chandra}, HST, and ATCA images for the quasar jet source PKS~2101--490. We extracted the radio to X-ray spectral energy distributions from seven regions of the 13$\\arcsec$ jet, and modeled the jet X-ray emission in terms of Doppler beamed inverse Compton scattering of the cosmic microwave background (IC/CMB) for a jet in a state of equipartition between magnetic and particle energy densities (\\S \\ref{sec:results}). Using this model, we derived a bulk Lorentz factor $\\Gamma \\sim 6$ and magnetic field strength of order 30~$\\mu$G.  A major goal of this work has been to assess the validity of the beamed, equipartition IC/CMB model for jet X-ray emission through the use of additional constraints: namely, independent estimates of jet power, and the location of an apparent cooling break in the synchrotron spectrum of one of the jet knots. Below we discuss the results of our analysis for each of these additional constraints.  The jet power predicted by the Doppler beamed, equipartition IC/CMB model was found to be in good agreement with independent order-of-magnitude estimates of jet power for this object, provided that $\\gamma_{\\rm min} \\sim 50$ in the jet, and the knots are significantly longer than the jet width, as implied by de-projection of the observed knot lengths (\\S \\ref{sec:jet_power}).  The brightest X-ray knot was detected in two HST filters, and the radio to optical data points were modelled as a broken power law with a standard cooling break $\\Delta \\alpha = 0.5$. The inferred break frequency was found to be consistent with the break frequency predicted using the IC/CMB model parameters along with the de-projected knot length that is implied by the small jet viewing angle required by the IC/CMB model (\\S \\ref{sec:2101_K6_broken_powerlaw}). However, we noted sub-structure in the F814W HST image of the bright optical knot, which is incompatible with a simple one-zone continuous injection model.  Finally, we note positive residuals consistent with a 6.4 keV Fe K$\\alpha$ line redshifted by the estimated z=1.04 reported by \\citet{marshall05}."
    },
    "1208/1208.5882_arXiv.txt": {
        "abstract": "{} {We study the enigmatic B[e] star MWC\\,300 to investigate its disk and binary with milli-arcsecond-scale angular resolution.} {We observed MWC\\,300 with the VLTI/AMBER instrument in the \\textit{H} and \\textit{K} bands and compared these observations with temperature-gradient models to derive model parameters. } {The measured low visibility values, wavelength dependence of the visibilities, and wavelength dependence of the closure phase directly suggest that MWC\\,300 consists of a resolved disk and a close binary. We present a model consisting of a binary and a temperature-gradient disk that is able to reproduce the visibilities, closure phases, and spectral energy distribution. This model allows us to constrain the projected binary separation ($\\sim$4.4\\,mas or $\\sim$7.9\\,AU), the flux ratio of the binary components ($\\sim$2.2), the disk temperature power-law index, and other parameters.} {} ",
        "introduction": "MWC\\,300 is an enigmatic early-type emission-line star whose nature and evolutionary state are not well known. \\citet{herbig60} classified MWC\\,300 as a pre-main-sequence star since nebulosity exists around the star. It is included in several catalogs of pre-main-sequence stars \\citep[e.g.,][]{the94}. However, MWC\\,300 was also listed as a probable B[e] supergiant \\citep[e.g.,][]{as76, appenzeller77, lamers98}. \\citet[][hereafter M04]{miroshnichenko04} carried out extensive spectroscopic and photometric studies on MWC\\,300, and strengthened the arguments in favor of the supergiant classification. M04 suggest that MWC\\,300 has an effective temperature of 19000\\,K, a distance of $1.8\\pm0.2$\\,kpc, and a luminosity of log\\,$L/L_{\\odot}=5.1\\pm0.1$.  Two-dimensional radiative transfer modeling shows that MWC\\,300 has a flared dusty disk rather than a spherical dusty envelope (M04). The modeled disk has an inner radius of $\\sim$11\\,mas and is viewed almost edge-on. \\citet{souza08} observed MWC\\,300 at 11.25\\,$\\mu$m using the BURST mode of VLT/VISIR. The obtained diffraction-limited  ($\\sim$0.3$''$) image shows that MWC\\,300 is partially resolved with an FWHM diameter of $69\\pm10$\\,mas. \\citet{monnier09} observed MWC\\,300 with the Keck telescope and measured a Gaussian FWHM diameter of $49\\pm3$\\,mas at 10.7\\,$\\mu$m.  MWC 300 exhibits radial velocity (RV) variations that suggest a binary \\citep{cl99}. \\citet{takami03} carried out spectro-astrometric observations and studied the absorption feature of H$\\alpha$. They detected a positional displacement of the line-emitting region of $17\\pm4$\\,mas with respect to the position of the continuum. Photospheric line RV variations detected by M04 also suggest a companion.   \\begin{table*} \\caption{Observation log of our VLTI/AMBER observation of MWC\\,300.} \\centering \\begin{tabular}{cccccccc} \\hline \\hline Date & Time & AT Configuration & Projected Baselines & PA & Seeing & DIT & Number of frames\\\\ & (UTC) & & (m) & $(^{\\circ})$ & ('') & (ms) & \\\\ \\hline 2009-04-18 & 07:37 - 07:46 & E0-G0-H0 & 13.90 / 27.78 / 41.68 & 66.2 & 0.5 & 200 & 2400\\\\ \\hline \\end{tabular} \\label{observation} \\end{table*}  \\begin{table*} \\caption{Scan ranges of the model parameters and their values for the best-fitting models with different inclination angles.} \\centering \\begin{tabular}{cccccccccc} \\hline \\hline & $i$\\footnotemark[1] & $\\theta_{\\rm{disk}}$\\footnotemark[2] & $r_{\\rm{in}}$\\footnotemark[3] & $r_{\\rm{out}}$\\footnotemark[4] & $T_{\\rm{in}}$\\footnotemark[5] & $\\alpha$\\footnotemark[6] & $r_{\\rm{s}}$\\footnotemark[7] & $\\kappa$\\footnotemark[8] & $\\chi^{2}_{\\rm{reduced}}$\\footnotemark[9] \\\\ & ($^{\\circ}$) & ($^{\\circ}$) & (mas) & (mas) & (K) & & (mas) &  & \\\\ \\hline Scan 1 & - & 0--180 & 1--50 & $r_{\\rm{in}}$--300 & 500--2000 & 0.0--1.0 & 1--1000 & 1.0--5.0 & -\\\\ Scan 2 & - & 0--10 & 10--20 & 100--300 & 800--1000 & 0.5--0.7 & 1--20 & 1.0--3.0 & -\\\\ Results & 80 & 4$\\pm3$ & 16.9$^{+1.8}_{-1.5}$ & 200$\\pm50$ & 890$\\pm20$ & 0.57$^{+0.3}_{-0.1}$ & 4.4$\\pm$0.2 & 2.2$^{+0.3}_{-0.2}$ & 2.47 \\\\ &  & & (30.4$^{+3.2}_{-2.7}$\\,AU) & (360$\\pm90$\\,AU) & & & (7.9$\\pm$0.4\\,AU) & & \\\\ \\hline Scan 1 & - & 0--180 & 1--50 & $r_{\\rm{in}}$--300 & 500--2000 & 0.0--1.0 & 1--1000 & 1.0--5.0 & -\\\\ Scan 2 & - & 10--30 & 5--15 & 100--300 & 800--1000 & 0.5--0.7 & 1--20 & 1.0--3.0 & -\\\\ Results & 70 & 17$\\pm5$ & 11.8$^{+1.0}_{-0.8}$ & 160$\\pm40$ & 900$^{+10}_{-20}$ & 0.58$^{+0.2}_{-0.1}$ & 4.4$\\pm$0.2 & 2.2$\\pm0.2$ & 2.45 \\\\ &  & & (21.2$^{+1.8}_{-1.4}$\\,AU) & (288$\\pm72$\\,AU) & & & (7.9$\\pm$0.4\\,AU) & & \\\\ \\hline Scan 1 & - & 0--180 & 1--50 & $r_{\\rm{in}}$--300 & 500--2000 & 0.0--1.0 & 1--1000 & 1.0--5.0 & -\\\\ Scan 2 & - & 20--40 & 5--15 & 100--300 & 800--1000 & 0.5--0.7 & 1--20 & 1.0--3.0 & -\\\\ Results & 60 & 29$\\pm7$ & 9.7$\\pm0.7$ & 140$\\pm40$ & 900$\\pm10$ & 0.58$\\pm0.1$ & 4.4$\\pm$0.2 & 2.2$^{+0.2}_{-0.1}$ & 2.44 \\\\ &  & & (17.5$\\pm1.3$\\,AU) & (252$\\pm72$\\,AU) & & & (7.9$\\pm$0.4\\,AU) & & \\\\ \\hline \\end{tabular}  \\footnotemark[1]{Disk inclination angle; }\\footnotemark[2]{PA of the long axis of the disk; }\\footnotemark[3]{inner disk radius; }\\footnotemark[4]{outer disk radius; }\\footnotemark[5]{disk temperature at the inner radius; }\\footnotemark[6]{temperature power-law index; }\\footnotemark[7]{projected binary separation along PA of $66.2^{\\circ}$; }\\footnotemark[8]{flux ratio between the stars; }\\footnotemark[9]{reduced $\\chi^{2}$ of the best-fitting model.} \\label{bestfittingparameter} \\end{table*}  Thanks to the development of near-infrared interferometry, we can now achieve a high spatial resolution of a few mas. It allows us to directly resolve the MWC\\,300 system and study the circumstellar disk structure, as well as the central binary. In this paper, we present the first near-infrared interferometric AMBER/VLTI observations of MWC\\,300. The observations are described in Section 2. In Section 3, we present a model consisting of a binary and a temperature-gradient disk that can reproduce the visibilities, closure phases, and SED. The results are discussed in Section 4, and conclusions are summarized in Section 5. ",
        "conclusions": "\\subsection{Disk inclination angle} \\label{diskinclination}  The emission line profile of MWC\\,300 and strong extinction along the line of sight suggest an almost edge-on disk (M04). The value of 80$^{\\circ}$ (i.e., almost edge-on) derived by  radiative transfer modeling (M04) is only a rough estimate, so we also investigated models with other inclination angles. The values of parameters for the best-fitting models with $i=60^{\\circ}$ and $70^{\\circ}$ are listed in Tabel \\ref{bestfittingparameter} (lines 4-9). These models can fit the observations with similar $\\chi^{2}$ values. Our results in Table \\ref{bestfittingparameter} show that some of our model parameters (e.g., $T_{\\rm{in}}$, $\\alpha$, $r_{\\rm{s}}$, and $\\kappa$) are independent of the disk inclination angle. However, $r_{\\rm{in}}$ and $\\theta_{\\rm{disk}}$ strongly depend on the disk inclination angle and cannot be constrained well. Our interferometric data do not yet allow us to constrain the inclination angle since our three visibilities were measured at the same PA. Therefore, more observations at different PAs are required in the future to improve the accuracy of the disk parameters.   \\subsection{Disk size} M04 modeled MWC\\,300 with a two-dimensional flared dusty-disk model and suggest that the inner boundary of the disk has a radius of $\\sim20$\\,AU and a temperature of $\\sim$1120\\,K. With the same assumed disk inclination angle of 80$^{\\circ}$, our best-fitting model suggests a slightly larger and cooler disk, with an inner radius of $\\sim$30\\,AU and a temperature at the inner radius of $\\sim$890\\,K. This radius is consistent with the dust sublimation radius of $\\sim$35\\,AU for 890\\,K and gray dust opacities, but is about 2.5 times larger than the 1500\\,K dust sublimation radius of $\\sim$12\\,AU \\citep{mm02}. This large inner disk radius and cool inner disk temperature may be explained by the existence of the secondary companion. Considering the minimum binary separation of 4.4\\,mas (7.9\\,AU), MWC\\,300 may have a circumbinary disk, and the secondary companion may be close enough to the disk inner radius to interact strongly with the disk \\citep{shi12}.    \\begin{table} \\caption{Flux contribution from different components at different wavelengths.} \\centering \\begin{tabular}{ccccc} \\hline \\hline & 1.6\\,$\\mu m$ & 1.9\\,$\\mu m$ & 2.1\\,$\\mu m$ & 2.4\\,$\\mu m$ \\\\ \\hline $f_{\\rm{primary}}$ & 0.491& 0.229& 0.124& 0.052\\\\ $f_{\\rm{secondary}}$ & 0.223& 0.104 & 0.056& 0.024\\\\ $f_{\\rm{disk}}$ & 0.286 & 0.667 & 0.82& 0.924\\\\ \\hline \\end{tabular} \\label{fluxratio} \\end{table}   \\subsection{Binarity and B[e] phenomenon}   Our interferometric data shows strong closure phases that suggest asymmetries. The asymmetries can be attributed to a binary or an extremely inhomogeneous disk. However, the observed closure phase signature is too strong to be caused by the asymmetry of the disk since the disk is only partially resolved \\citep{monnier06}. The observed closure phase is therefore the result of a secondary companion.  Adopting the distance of 1.8\\,kpc, we obtain a small projected binary separation of $\\sim$7.9\\,AU. However, due to our one-dimensional $uv$ coverage, we can only derive the projected separation of the binary along PA of $66.2^{\\circ}$ and not its exact separation. The lower-right hand panel of Fig. \\ref{eomodel} indicates the possible locations of the secondary companion. All companion positions on the green line correspond to a projected binary separation of 7.9\\,AU. The figure shows that MWC\\,300 may have a circumbinary disk if the binary separation is small enough \\citep[as, for example, in HD\\,62623; ][]{plets95,millour09}. However, considering the AMBER AT's field of view, the binary separation can be large enough (with a maximum value of 150\\,mas) for the disk to be a circumprimary disk \\citep[as, for example, in HD\\,87643; ][]{millour09}. More interferometric observations are required in order to constrain the exact separation and PA of the binary.  Several other B[e] stars were also found to be binaries: $\\eta$ Car, MWC\\,349A, HD\\,87643, V921\\,Sco, and HD\\,327083. In HD\\,87643 \\citep{millour09} and V921\\,Sco \\citep{kraus12a}, circumprimary disks were resolved, whereas in HD\\,327083 \\citep{wheelwright12} a circumbinary disk was detected. The $\\eta$ Car binary was found by photometric observations \\citep{damineli97}. The binarity of MWC\\,349A has not been directly confirmed, but was suggested by \\citet{hofmann02}. Although the exact role of binarity in the B[e] phenomenon is still an open question, it is possible that the disks are the result of mass exchange episodes between the stellar components of the B[e] stars \\citep{sheikina00, miroshnichenko07a, miroshnichenko07b, miroshnichenko09, kraus10, kraus12b}. The slowing down of the radiative wind by the companion gravitational effect or wind shock may also be an explanation \\citep{plets95}. Future interferometric observations will allow us to improve our understanding, since these observations will be able to determine the orbits of the B[e] binaries and to investigate the interaction of the binary companion with the disk or the wind."
    },
    "1208/1208.2000_arXiv.txt": {
        "abstract": "During its 16 years of service the Rossi X-ray Timing Explorer (\\xte )  mission has provided an extensive  archive of data, which will serve as a primary source of high cadence observations of variable X-ray sources for fast timing studies. It is, therefore, very important to have the most reliable calibration of \\xte\\ instruments.   The Proportional Counter Array (PCA) is the primary instrument on-board \\xte\\ which provides data in 3-50 keV energy range with sub-millisecond time resolution in up to 256 energy channels. In 2009 the RXTE team revised the response residual minimization method used to derive the parameters of the PCA physical model. The procedure is based on the residual minimization between the  model spectrum for Crab nebula emission and a calibration data set consisting of a number of spectra from the Crab and the on-board Am$_{241}$ calibration source, uniformly covering the whole \\xte\\ mission operation period. The new method led to a much more effective model convergence and  allowed for better understanding of the PCA energy-to-channel relationship. It greatly improved the response matrix performance. We describe the new version of the \\xte/PCA response generator PCARMF v11.7 (HEASOFT Release 6.7)along with the corresponding energy-to-channel conversion table (verson {\\it e05v04}) and their difference from the previous releases of PCA calibration. The new PCA response adequately represents the spectrum of the calibration sources and successfully predicts the energy of the narrow iron emission line in Cas-A throughout the RXTE mission. ",
        "introduction": "The Rossi X-ray Timing Explorer (\\xte) was launched on December 30, 1995 and successfully operated until January 4, 2012.  \\xte\\ is an X-ray observatory with a powerful and unique combination of large collecting area, broad-band spectral coverage, high time resolution, flexible scheduling, and ability of quick response and frequent monitoring of time-critical targets of opportunity. \\xte\\ observations have led to breakthroughs in our understanding of physics of strong gravity, high density, and intense magnetic field environments found in neutron stars, galactic and extragalactic black holes and other sources. The mission combined two pointing instruments, the Proportional Counter Array (PCA) developed to provide data for energies between 3 and 50 keV, and the High Energy X-ray Timing Experiment \\citep[HEXTE;][]{hexte}  covering the 20-250 keV energy range. These instruments were equipped with collimators yielding a FWHM of one degree. In addition, RXTE carried an All-Sky Monitor \\citep[ASM;][]{asm} that scans about 80\\% of the sky every orbit, allowing monitoring at time scales of 90 minutes or longer. Data from PCA and ASM are processed on board by the Experiment Data System (EDS).  The PCA is array of five large-area proportional counter units (PCUs) designed to perform observations of bright X-ray sources with high timing and modest spectral resolution. The main chamber of each PCU is divided into three volumes or layers filed with xenon. In addition, all PCUs were initially equipped with a propane-filled '\"veto\" layer in front of the top xenon layer.  The calibration of the PCA, as well as the details of its design and operation, are described in \\citet[][J06 hereafter]{jah06}. The response generation software for the PCA is based on physical model of the instrument. The main components of the model are the quantum efficiency, which gives the probability of an X-ray  photon to be absorbed in one of the detector volumes, and the redistribution matrix, which provides the probability for a photon to be detected in one of the PCU energy channels.  The  model has a complex dependence on many parameters, which have to be properly optimized to minimize a difference between the predicted model and the observed spectrum of one or more calibration sources, i.e. sources with well known spectral characteristics. Implementation of an effective parameter optimization procedure is vital in performing this task.  The set of PCA  parameters describing the instrument response since the start of the mission and until 2004 has been calculated in J06. However, calibration observations of the Crab and other sources after 2004 suggested that the model and its parameters have to be updated to provide a consistent response for new science observations. In 2009 we revised the PCA model and the response  minimization method. The new model provided  significant response improvement for the entire mission span, and especially for the data collected after 2004.  In this paper we describe in detail the new fitting procedure and the results of the PCA response modeling. The paper is structured as follows. In the next section we provide a brief review of the PCA physical model and the response generation software. The calibration data is described in \\S 3. In \\S \\ref{setup} we provide the details of our PCA model implementation and the response minimization fitting environment in XSPEC astrophysical fitting package. We describe and discuss the results in \\S \\ref{disc}. Conclusions follow in \\S \\ref{summary}. ",
        "conclusions": "\\label{summary}  We present a new \\xte\\/PCA response version v11.7. This response is based on data set presenting the entire \\xte\\ mission span. While the new response is largely based on the physical model developed in J06, some significant modifications are made, especially for energy-to-channel conversion relationship. The new response shows a superior performance with respect to the previous \\xte\\ response versions."
    },
    "1208/1208.0005_arXiv.txt": {
        "abstract": "We apply two different algorithms to search for mass segregation to a recent observational census of the $\\rho$~Ophiuchi star forming region. Firstly, we apply the $\\Lambda_{\\rm MSR}$ method, which compares the minimum spanning tree (MST) of a chosen subset of stars to MSTs of random subsets of stars in the cluster, and determine the mass segregation ratio, $\\Lambda_{\\rm MSR}$. Secondly, we apply the $m-\\Sigma$ method, which calculates the local stellar surface density around each star and determines the statistical significance of the average surface density for a chosen mass bin, compared to the average surface density in the whole cluster. Using both methods,  we find no indication of mass segregation (normal or inverse) in the spatial distribution of stars and brown dwarfs in $\\rho$~Ophiuchi. Although $\\rho$~Ophiuchi suffers from high visual extinction, we show that a significant mass segregation signature would be detectable, albeit slightly diluted, despite dust obscuration of centrally located massive stars. ",
        "introduction": "Most stars form in groups, clusters, and larger associations. In order to understand the star formation process, it is desirable to quantify the spatial distribution of stars in different star forming regions, so that a clear picture of the formation and evolution of each region can be drawn. It is possible to measure the amount of substructure in a region \\citep[e.g.\\,\\,by using the $\\mathcal{Q}$-parameter,][]{Cartwright04} and to quantify the amount of mass segregation (e.g.\\,\\,the $\\Lambda_{\\rm MSR}$ method, \\citealp{Allison09a}, or the $m-\\Sigma$ method, \\citealp{Maschberger11}). Additionally, statistical methods can be applied to find clusters against a background field \\citep[e.g.][]{Gutermuth09,Schmeja11}.  \\citet{Allison09a} found that the amount of mass segregation in the ONC could be quantified by comparing the minimum spanning trees (MSTs) of chosen subsets of stars to the MSTs of random sets of stars. If the MST of the most massive stars is shorter than the MSTs of random subsets of cluster stars, then the cluster is mass segregated. The ONC is mass segregated \\citep[see also][]{Hillenbrand98}, and the same signature was found by \\citet{Sana10} in Trumpler~14.  However, \\citet{Parker11b} found that the most massive stars in the Taurus association were `inversely mass segregated', i.e.\\,\\,anti-clustered with respect to randomly chosen stars. Is mass segregation therefore a dynamical process \\citep[as postulated by][]{Allison09b}, rather than a primordial outcome of star formation (in hydrodynamical simulations of star cluster formation, primordial mass segregation occurs as part of the competitive accretion process, e.g.\\,\\,\\citealp{Maschberger10}, \\citealp{Maschberger11})? To answer this question, we must first search for mass segregation in other young star forming regions, ideally using independent methods.  \\begin{figure*} \\begin{center} \\rotatebox{270}{\\includegraphics[scale=0.65]{Oph_map_wir.ps}} \\end{center} \\caption[bf]{A map of $\\rho$~Ophiuchi showing the 255 objects in our dataset (restricted to the WIRCam field). The 20 least massive cluster members (masses up to 0.03\\,M$_\\odot$) are shown by the blue crosses and the 20 most massive cluster members (masses down to 1.63\\,M$_\\odot$) are shown by the large red dots. The solid lines indicate the extent of the WIRCam field.} \\label{Oph_map} \\end{figure*}  In this paper we search for mass segregation in $\\rho$~Ophiuchi. This cluster suffers heavily from differential extinction, so an accurate and self-consistent determination of stellar masses is difficult. However, recent spectroscopic surveys \\citep{Alves12,Erickson2011,Geers2011,Muzic2011} have probed the low-mass end of the IMF and allowed a complete census of the cluster to be made. We describe the observational sample in Section~\\ref{observations}, we describe the methods used to quantify mass segregation in Section~\\ref{method}, we present our results in Section~\\ref{results}, we discuss the results and the potential effects of extinction in Section~\\ref{discuss} and we conclude in Section~\\ref{conclude}. ",
        "conclusions": "\\label{conclude}   We have used an observational census of $\\rho$~Ophiuchi, which was recently enhanced by several surveys probing the substellar domain of the IMF, to search for possible mass segregation signatures in the spatial distribution of stars and brown dwarfs in this cluster.  We have utilised two different algorithms. Firstly, we used the $\\Lambda_{\\rm MSR}$ technique \\citep{Allison09a}, which compares the minimum spanning tree (MST) of a chosen subset of stars, to the MSTs of randomly chosen stars in the cluster. If the MST length of a chosen subset is shorter than the MST length of the random objects, then the cluster is mass segregated.  Secondly, we have used the $m-\\Sigma$ plot, which compares the local surface density surrounding massive stars to the the average surface density of all of the stars in the cluster. By this definition, a cluster is mass segregated if the massive stars have a significantly higher than average surface density. Our conclusions are as follows:  (i) The $\\Lambda_{\\rm MSR}$ technique finds that the most massive stars show hints of being inversely mass segregated, with $\\Lambda_{\\rm MSR} = 0.89^{+0.09}_{-0.13}$ for the 20 most massive stars. However, $\\Lambda_{\\rm MSR}$ is consistent with there being no mass segregation of the 26 most massive stars, and so on.  The least massive stars show no clear deviation from $\\Lambda_{\\rm MSR} = 1$. \\newline (ii) The $m-\\Sigma$ distribution also suggests that the most massive stars may be inversely mass segregated (but with no strong statistical significance), and with no difference in the distribution of low-mass stars compared to the cluster average. \\newline (iii) The high levels of extinction in $\\rho$~Oph may mean that some members are missing from the dataset. However, we have demonstrated that a significant difference in the spatial distribution of a group of objects would still be found by both the  $\\Lambda_{\\rm MSR}$ and $m-\\Sigma$ methods.\\newline  In order to understand the star formation process in different clusters, we suggest applying both mass segregation algorithms in tandem to build up a census of the spatial distribution of stars in different star forming regions."
    },
    "1208/1208.0011_arXiv.txt": {
        "abstract": "{In this contribution I present a review of star formation in clusters. I begin by discussing the various definitions of what constitutes a star cluster, and then compare the outcome of star formation (IMF, multiplicity, mass segregation and structure and morphology) in different star-forming regions. I also review recent numerical models of star formation in clusters, before ending with a summary of the potential effects of dynamical evolution in star clusters.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0820_arXiv.txt": {
        "abstract": "$\\mu\\,$Columbae is a prototypical weak-wind O-star for which we have obtained a high-resolution X-ray spectrum with the \\chan\\ \\letg/\\acis instrument and a low resolution spectrum with {\\em \\suz}.  This allows us, for the first time, to investigate the role of X-rays on the wind structure in a {\\it bona fide} weak-wind system and to determine whether there actually is a massive, hot wind.  The X-ray emission measure indicates that the outflow is an order of magnitude greater than that derived from UV lines and is commensurate with the nominal wind-luminosity relationship for O-stars.  Therefore, the ``weak-wind problem''---identified from cool wind UV/optical spectra---is largely resolved by accounting for the hot wind seen in X-rays.  From X-ray line profiles, Doppler shifts, and relative strengths, we find that this weak-wind star is typical of other late O dwarfs.  The X-ray spectra do not suggest a magnetically confined plasma---the spectrum is soft and lines are broadened; \\suz spectra confirm the lack of emission above 2 keV. Nor do the relative line shifts and widths suggest any wind decoupling by ions.  The He-like triplets indicate that the bulk of the X-ray emission is formed rather close to the star, within 5 stellar radii. Our results challenge the idea that some OB stars are ``weak-wind'' stars that deviate from the standard wind-luminosity relationship. The wind is not weak, but it is hot and its bulk is only detectable in X-rays. ",
        "introduction": "The outflow of stellar winds from massive OB-type stars is an important process which affects both the chemical enrichment and kinetics of the interstellar medium \\citep[e.g.][]{Leitherer:Robert:Drissen:1992}.  The mass-loss itself is enough to change the evolution of the star, which ends its life in a supernova explosion, also profoundly changing its environment. Hence, quantitative understanding massive star winds is important not only as a basic component of stellar astrophysics, but also for understanding cosmic feedback on galactic scales throughout cosmic history.  While some basic physics of stellar winds in massive stars is well established---that winds are accelerated by photoelectric absorption of the intense ultraviolet radiation field by a multitude of metal lines, that an instability can lead to wind-shocks which generate X-rays---there are still puzzles to be solved.  One of these is the ``weak-wind'' problem, in which UV line diagnostics clearly show a wind signature in classical P~Cygni line profiles, but modeled mass loss rates can be discrepant by more than an order of magnitude from values expected based on O-star statistical trends and theoretical foundations, specifically the wind momentum-luminosity relation \\citep[e.g.][]{Puls:Kudritzki:Herrero:al:1996}.  Factors of a few in mass loss are enough to be significant for stellar evolution and cosmic feedback \\citep[e.g.][]{Puls:Vink:Najarro:2008}.  It has long been known that photoionization by X-rays can alter the ionization balance in the wind regions where the UV lines are formed \\citep{Waldron:1984,Macfarlane:Waldron:al:1993}.  There is a theoretical degeneracy in that different values of mass loss rate ($\\dot M$) and X-ray luminosity ($L_x$) can produce very similar UV line profiles \\citep{Puls:Vink:Najarro:2008, Marcolino:Bouret:al:2009}; direct knowledge of the X-ray spectrum is thus important for reliable determination of wind parameters.  Another possibility is that cool and hot plasma emission originate from different volumes or densities; that is, clumping can affect the interpretation \\citep{Hamann:Feldmeier:Oskinova:2008}.  \\mucol belongs to the weak-wind domain defined by \\citet{Lucy:2010} in which a star's rate of mechanical energy loss in a radiatively driven wind is less than the radiative output from nuclear burning; Lucy showed that there is a huge disparity between the theoretically expected $\\dot M$ and values derived from UV and optical spectra. \\citet{Lucy:2012} developed a phenomenological model and suggested that in low-luminosity O-type stars, the volumetric roles of hot and cool gas are possibly reversed compared to O-type stars of high luminosity; thus in the weak-wind stars, a larger volume is occupied by the hot gas than by the cool gas.  \\mucol (HD~38666) is an O9.5~V runaway (and single) star and is one of the weakest wind Galactic O-type stars \\citep[e.g., see Figs.\\ 39 and 41 in][]{Martins:Schaerer:al:2005}.  These factors are what motivated our spectroscopic study of the prototypical weak-wind system, \\mucol, at high-resolution with \\chan and at low resolution but greater sensitivity at higher energies with \\suz.  In this letter, we concentrate on the primary empirical results from the X-ray spectral analysis of \\mucol.  A subsequent paper will investigate the influence of the X-rays on the cool wind component (Todt et. al, in preparation).  \\begin{figure*}[!htb] \\centering\\leavevmode \\includegraphics[width=0.98\\textwidth,viewport=1 1 530 140]{mucol_count_rate_spec-04_color.ps} \\caption{\\mucol\\ \\chan/\\letg/\\acis spectrum (left); black: observed count-rate; thin-gray line (red in the online version): model; the lower panel shows the $\\chi^2$ residuals against a broken powerlaw APEC model modified for photoexcitation of triplets.  Lines are broader than the instrumental width. The \\suz spectrum is shown on the right (black).  For conciseness, we have summed the counts from the three detectors; such is not recommended when fitting, but it provides a good summary visualization of the data.  The folded \\chan-derived model is in gray (red in the online version), and residuals below.  We emphasize that the \\chan-derived model was not fit to the \\suz spectrum, only folded through the response to provide model counts and residuals; the model agrees very well without any adjustments.  It is significant that there is little or no flux detected above $2\\kev$ where \\suz has substantial sensitivity.} \\label{fig:mucolspec} \\end{figure*}  \\begin{deluxetable}{ll|ll} \\tablecolumns{4} \\tablewidth{0.95\\columnwidth} \\tablecaption{\\mucol Properties} \\tablehead{ \\colhead{Property}&  \\colhead{Value} &\\colhead{Property}&  \\colhead{Value} } \\startdata Spectral Type& O9.5 V& $d\\,$[pc]\\tablenotemark{a}& 408\\\\  $T_\\mathrm{eff}\\,$[K]\\tablenotemark{b}& $33000$& $R/R_\\odot$\\tablenotemark{b,e}& $4.6$\\\\  $\\log(L_\\mathrm{bol}/L_\\odot)$\\tablenotemark{b,e}& 4.4& $v_\\infty[\\kms]$\\tablenotemark{b}& $1200$\\\\  $\\dot M_{uv}\\,\\mathrm{[M_\\odot/yr]}$\\tablenotemark{b}& $10^{-9.5}$& $N_\\mathrm{H}\\,[\\mathrm{cm^{-2}}]$\\tablenotemark{c}& $5\\times10^{19}$\\\\  $f_x(1-40\\mang)[\\mathrm{cgs}]$\\tablenotemark{d,f}& $6\\times10^{-13}$& $\\log(L_\\mathrm{x}/L_\\mathrm{bol})$\\tablenotemark{d}& $-6.9$\\\\ $EM_x [\\cmmthree]$\\tablenotemark{d}& $10^{54.1}$& $T_\\mathrm{max}[\\mk]$\\tablenotemark{d}& $4.4$\\\\ $v_{\\infty,x}[\\kms]$\\tablenotemark{d}& $1600\\,(\\pm275)$% &&\\\\ $f/i(${\\eli{O}{7}}$)$\\tablenotemark{g}& $<0.01$ & $r($\\eli{O}{7}$)$\\tablenotemark{g}& $<3.3$\\\\ $f/i(${\\eli{Ne}{9}}$)$\\tablenotemark{g}& $0.04$--$0.14$ & $r($\\eli{Ne}{9}$)$\\tablenotemark{g}& $2.3$--$4.4$\\\\ $f/i(${\\eli{Mg}{11}}$)$\\tablenotemark{g}& $>0.2$ & $r($\\eli{Mg}{11}$)$\\tablenotemark{g}& $>2.1$ \\enddata \\tablenotetext{a}{From the Hipparcos parallax, as re-evaluated by \\citet{vanLeeuwen:2007}} \\tablenotetext{b}{\\citet{Martins:Schaerer:al:2005}} \\tablenotetext{c}{\\citet{Cassinelli:al:1994,Howk:Savage:Fabian:1999}} \\tablenotetext{d}{This work} \\tablenotetext{e}{Adjusted for the adopted distance.} \\tablenotetext{f}{The flux is as observed at Earth, with foreground absorption.} \\tablenotetext{g}{$f/i$ gives the ratio of the forbidden to intercombination line fluxes, and $r$ is the radius of formation in units of the stellar radius, as derived from PoWR models.} \\label{tbl:props} \\end{deluxetable} ",
        "conclusions": "Our results challenge the idea that some OB stars are ``weak-wind'' stars that deviate from the standard wind-luminosity relationship. From high-resolution X-ray spectrum of \\mucol, specifically He-like lines and the total emission measure, this star does not appear unusual relative to other O-stars except for its weak-wind status.  Its X-ray emission measure, line widths, and centroids are in good agreement with the OB main sequence star results of \\citet{Waldron:Cassinelli:2007}.  The volume emission measure of the X-ray emitting plasma must be very much larger than the cool, UV-emitting plasma, and we believe that the weak-wind problem is reduced or eliminated when the hot and dominant component of the wind is taken into account.  The wind is not weak, but it is hot and its bulk is only detectable in X-rays."
    },
    "1208/1208.0966_arXiv.txt": {
        "abstract": "A 2D linear theory of the instability of Sweet-Parker (SP) current sheets is developed in the framework of Reduced MHD. A local analysis is performed taking into account the dependence of a generic equilibrium profile on the outflow coordinate. The plasmoid instability [Loureiro \\etal, Phys. Plasmas {\\bf 14}, 100703 (2007)] is recovered, i.e., current sheets are unstable to the formation of a large-wave-number chain of plasmoids ($k_{\\rm max}\\Lsheet \\sim S^{3/8}$, where $k_{\\rm max}$ is the wave-number of fastest growing mode, $S=\\Lsheet V_A/\\eta$ is the Lundquist number, $\\Lsheet$ is the length of the sheet, $V_A$ is the Alfv\\'en speed and $\\eta$ is the plasma resistivity), which grows super-Alfv\\'enically fast ($\\gmax\\tau_A\\sim S^{1/4}$, where $\\gmax$ is the maximum growth rate, and $\\tau_A=\\Lsheet/V_A$). For typical background profiles, the growth rate and the wave-number are found to {\\it increase} in the outflow direction. This is due to the presence of another mode, the Kelvin-Helmholtz (KH) instability, which is triggered at the periphery of the layer, where the outflow velocity  exceeds the Alfv\\'en speed associated with the upstream magnetic field. The KH instability grows even faster than the plasmoid instability, $\\gmax \\tau_A \\sim  k_{\\rm max} \\Lsheet\\sim S^{1/2}$. The effect of viscosity ($\\nu$) on the plasmoid instability is also addressed. In the limit of large magnetic  Prandtl numbers, $Pm=\\nu/\\eta$, it is found that $\\gmax\\sim S^{1/4}Pm^{-5/8}$ and $k_{\\rm max} \\Lsheet\\sim S^{3/8}Pm^{-3/16}$, leading to the prediction that the critical Lundquist number for plasmoid instability in the $Pm\\gg1$ regime is $\\Scrit\\sim 10^4Pm^{1/2}$. These results are verified via direct numerical simulation of the linearized equations, using a new, analytical 2D SP equilibrium solution. ",
        "introduction": "\\label{intro} Magnetic reconnection~\\cite{biskamp_magnetic_2000,priest_magnetic_2000, zweibel_magnetic_2009, yamada_magnetic_2010} is a ubiquitous plasma physics phenomenon, characterized by the rapid reconfiguration of the magnetic-field topology. Solar flares~\\cite{shibata_solar_2011} and magnetospheric substorms~\\cite{oieroset_situ_2001} are two prominent examples of events where reconnection plays a key role. Plasma dynamics in many laboratory experiments is also critically determined by magnetic reconnection; examples are the sawtooth~\\cite{hastie_sawtooth_1997} and the tearing instabilities~\\cite{furth_finite-resistivity_1963,rutherford_nonlinear_1973} in magnetic-confinement-fusion devices, or the reconnection of high-energy-density, laser-produced plasma bubbles~\\cite{nilson_magnetic_2006,willingale_proton_2010, dong_plasmoid_2012}.  Along with fast reconnection rates, many observations~\\cite{lin_plasmoids_2008} of magnetic reconnection phenomena display one intriguing feature: the formation, and subsequent ejection from the current sheet, of coherent secondary structures, often referred to as plasmoids (also known as blobs, flux ropes or secondary magnetic islands). There is abundant direct evidence for the presence of these structures in the Earth's magnetotail~\\cite{zong_cluster_2004, eastwood_observations_2005,chen_observation_2008} and in solar flares~\\cite{lin_direct_2005,ciaravella_current_2008, bemporad_spectroscopic_2008,lin_investigation_2009, nishizuka_multiple_2010, takasao_simultaneous_2012}. In magnetic confinement fusion devices, plasmoid generation seems to be less certain, though there are reports of the observation of secondary magnetic structures correlated to $m/n=1/1$ and $m/n=2/1$ magnetic islands on the TEXTOR~\\cite{donne_evidence_2005,liang_observations_2007} and JET~\\cite{salzedas_secondary_2011} tokamaks. On TEXTOR, high-resolution measurements of electron temperature fluctuations show structures which hint at plasmoid formation during sawtooth crashes~\\cite{park_observation_2006,park_comparison_2006, munsat_letter:_2007}. Finally, recent laser-plasma experiments where reconnection is conjectured to occur also show evidence for plasmoid formation~\\cite{willingale_proton_2010,dong_plasmoid_2012}.  Direct numerical simulations of reconnection processes concur with observations in displaying ubiquitous evidence for plasmoid formation. Plasmoids have been reported in numerical simulations using various physical models, ranging from kinetic~\\cite{daughton_fully_2006,drake_formation_2006, karimabadi_multi-scale_2007,daughton_transition_2009, daughton_influence_2009} to Hall-MHD~\\cite{shepherd_comparison_2010,huang_onset_2011} and to single fluid MHD~\\cite{park_reconnection_1984,steinolfson_nonlinear_1984,lee_multiple_1986, biskamp_magnetic_1986,jin_twodimensional_1991,loureiro_x-point_2005-1, lapenta_self-feeding_2008,samtaney_formation_2009,loureiro_turbulent_2009, bhattacharjee_fast_2009, cassak_scaling_2009-1, huang_scaling_2010, skender_instability_2010, loureiro_magnetic_2012}. Plasmoid formation has also been reported in numerical simulations of reconnection in relativistic plasmas, both resistive~\\cite{komissarov_tearing_2007} and kinetic~\\cite{zenitani_role_2008,liu_particle_2011,sironi_acceleration_2011,cerutti_beaming_2012}. Numerical studies tailored to address specific reconnection contexts such as the solar corona~\\cite{riley_``bursty_2007,barta_dynamics_2008, bettarini_spontaneous_2010,barta_spontaneous_2010}, the Earth's magnetotail~\\cite{karimabadi_magnetic_1999, jin_2.5_2001}, magnetic young stellar objects~\\cite{goodson_jets_1999,goodson_jets_1999-1,uzdensky_magnetic_2004}, fusion experiments~\\cite{park_reconnection_1984,biskamp_magnetic_1986} and laser-plasma interactions~\\cite{fox_fast_2011}, though different from each other in a number of details, again all appear to agree on the basic fact that plasmoids are generated in reconnecting current sheets.  The plasmoid dynamics inferred from observations and seen in numerical simulations strongly suggest the very \\textit{opposite} of the laminar, steady-state reconnection scenarios that have dominated the field for much of its history [the  Sweet-Parker (SP)~\\cite{parker_sweets_1957,sweet_neutral_1958} and the Petschek~\\cite{petschek_magnetic_1964} models, and, more recently, the Hall reconnection paradigm~\\cite{birn_geospace_2001}]. Magnetic reconnection in the presence of plasmoids appears to be a highly time-dependent, bursty process, which can only be described in a statistical manner~\\cite{shibata_plasmoid-induced-reconnection_2001,barta_spontaneous_2010, fermo_statistical_2010,uzdensky_fast_2010, fermo_comparison_2011, loureiro_magnetic_2012}. Furthermore, in addition to their key role in setting the reconnection rate in both laminar~\\cite{shibata_plasmoid-induced-reconnection_2001, lapenta_self-feeding_2008, daughton_transition_2009, daughton_influence_2009,bhattacharjee_fast_2009,cassak_scaling_2009-1,huang_scaling_2010, uzdensky_fast_2010,huang_onset_2011,loureiro_magnetic_2012} and turbulent~\\cite{loureiro_turbulent_2009,skender_instability_2010} plasmas, there is numerical and observational evidence that plasmoids may be critical in explaining electron acceleration in reconnection sites~\\cite{drake_electron_2006, chen_observation_2008, oka_island_2010, oka_electron_2010}.  In a previous paper~\\cite{loureiro_instability_2007} (henceforth referred to as Paper I) we attempted to understand the origin of plasmoid formation in reconnection sites by analysing the linear stability of large-aspect-ratio, SP current sheets. These were found to be violently unstable to the formation of plasmoid chains, the  fastest growing wave number scaling as $k_{\\rm max} \\Lsheet \\sim S^{3/8}$, with corresponding growth rate $\\gmax \\tau_A\\sim S^{1/4}$, where $\\Lsheet$ is the length of the current layer, $\\tau_A=\\Lsheet/V_A$ is the Alfv\\'en time ($V_A$ is the Alfv\\'en speed) and $S$ is the Lundquist number, $S=\\Lsheet V_A/\\eta$, where $\\eta$ is the magnetic diffusivity. Since $S\\gg 1$ in most applications of interest, this theory predicts the formation of multidinous plasmoids growing super-Alfv\\'enically; the immediate implication is that stable reconnecting current sheets at large values of the Lundquist number cannot exist. These results have since been confirmed in direct numerical simulations~\\cite{samtaney_formation_2009, ni_linear_2010}, and extended to account for the effect of a finite component of the magnetic field perpendicular to the reconnection plane~\\cite{baalrud_reduced_2012} and into the two-fluid regime~\\cite{baalrud_hall_2011}.  The analysis of Paper I considered a very simplified model background equilibrium, intended to retain only what we viewed as the most important features of a SP current sheet: a reconnecting magnetic field, $\\bm B_{eq}=(0,B_y(x))$ ($x$ is the inflow direction, $y$ the outflow direction) and an incompressible flow defined by the stream function $\\phi_{eq}=\\Gamma_0xy$, where $\\Gamma_0=V_A/\\Lsheet$ is the flow shearing rate. The analytical derivation in Paper I did not, therefore, take into account potentially important effects, such as the variation of the reconnecting magnetic field and of the outflow speed along the layer (i.e., along the $y$-direction in our chosen geometry), or the reconnected magnetic field.  In this work, we generalize the results of Paper I to a more realistic, two-dimensional model of the current sheet. Using an approach in the spirit of WKB theory (justified by the expectation that the most unstable wave-number will be very large, $k_{\\rm max}\\Lsheet\\gg1$), we derive the dispersion relation for the plasmoid instability as a slow function of the position along the sheet, $y_0$. We find that the scalings of the maximum growth rate ($\\gmax$) and wave-number ($k_{\\rm max}$) derived in Paper I hold true in a central, finite-sized patch of  the current sheet; however, the growth rate and wave number are now parametrized nontrivially by $y_0$. Surprisingly, we also discover that for a generic background equilibrium configuration, the maximum growth rate and wave-number of the instability \\textit{increase} with $y_0$ (i.e., outwards). As we show in this paper, a special point exists, $y_{0,\\rm crit}$, beyond which the assumptions invoked in our calculation break down. This is the Alfv\\'en Mach point of the system, where the magnitude of the outflow velocity (an increasing function of $y_0$) becomes equal to the value of the Alfv\\'en speed based on the upstream magnetic field (a decreasing function of $y_0$). Beyond that point, the current sheet becomes unstable to a different mode: the Kelvin-Helmholtz (KH) instability, whose growth rate and wave-number dependence we also derive analytically.  The other main result of this paper is the study of the effect of a large viscosity $\\nu$ (parametrized by the magnetic Prandtl number $Pm = \\nu/\\eta\\gg 1$) on the plasmoid instability. The large-Prandtl-number regime is pertinent to various astrophysical applications, e.g., the interstellar medium~\\cite{kulsrud_spectrum_1992}, and to fusion plasmas~\\cite{park_reconnection_1984}, and so it is important to understand how large $Pm$ affects plasmoid formation and dynamics.  Our analytical results are complemented with a direct numerical solution of the full set of linearized equations.   This paper is organized as follows. In~\\secref{sec:heuristic}, we present a heuristic derivation of our main results. A more rigorous approach to the problem begins in \\secref{sec:prob_setup}, where the equations to be solved are laid out and the expected properties of the background equilibrium are discussed (a more quantitative discussion of the constraints that the background equilibrium should satisfy can be found in Appendix~\\ref{equilibrium}, where an analytical 2D SP-like current-sheet equilibrium is obtained). The core of the analytical calculation is presented in \\secref{sec:lin_theory}. The KH instability of the current sheet is derived in~\\secref{KH}. Results of the direct numerical solution of the linear equations are presented in~\\secref{numerics}. The effect of viscosity on the instability is addressed in~\\secref{sec:viscosity}. Finally, a discussion of the results and conclusions can be found in~\\secref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} In this paper, a two-dimensional linear theory of the instability of large-aspect-ratio, Sweet-Parker-like current sheets is presented. This work is a direct generalization of our previous results~\\cite{loureiro_instability_2007} (Paper I), where the simple equilibrium used was only a good model of a current sheet in the immediate vicinity of $y=0$ ($y$ is the outflow direction).  In the work presented here a general 2D SP-like current sheet equilibrium is considered. As in Paper I, we conclude that large-aspect-ratio Sweet--Parker current sheets are violently unstable to high-wave-number tearing-like perturbations, and the same scalings of the growth rate with the Lundquist number $S=L V_A/\\eta$ are obtained here: $\\gmax \\tau_A\\sim S^{1/4}$ and $k_{\\rm max} \\Lsheet \\sim S^{3/8}$ --- see \\eqs{k_max}{gamma_max}. The plasmoid chain is formed inside a boundary layer whose width scales as $\\deltain/\\deltacs\\sim S^{-1/8}$. These scalings have been confirmed via direct numerical simulation~\\cite{samtaney_formation_2009,ni_linear_2010}.  The more general approach employed in this paper has allowed us to calculate the growth rate of the plasmoid instability as a function of the position along the current sheet, $y_0$. The dependence of $\\gmax$ and $k_{\\rm max}$ on $y_0$ is a nontrivial function of the particular equilibrium considered and, in the absence of a known analytical solution to the SP problem, cannot be evaluated explicitly. However, for $y_0/\\Lsheet\\ll1$ we make use of the semi-analytical results of Uzdensky and Kulsrud~\\cite{uzdensky_viscous_1998} and present an exact solution --- Eqs.~(\\ref{k_max_uz},~\\ref{gamma_max_uz}). The most unstable wave-number and corresponding growth rate are then found to \\textit{increase} with distance from the center. Under general conditions (Syrovatskii-like upstream magnetic-field profile and outflow profile increasing monotonically along the layer), we show that the same result holds true at arbitrary $y_0/\\Lsheet\\sim1$. This finding is somewhat counterintuitive: \\textit{a priori}, one could expect that the increasing strength of the reconnected field along the sheet, as well as the shear in the ouflow (in the $y$ direction), would provide a stabilising effect. Our calculation shows, however, that both are irrelevant to the instability. An intuitive understanding of why that should be so can be gained by comparing the strength of the upstream and the downstream magnetic fields at the boundary of the inner (plasmoid) layer, $\\xi=\\deltain$: \\be \\left. \\frac{B_y}{B_x}\\right|_{x\\sim\\deltain}\\sim  \\frac{\\deltain}{\\deltacs} S^{1/2} \\sim S^{3/8}\\gg 1, \\ee i.e., even at the scale of the inner layer the reconnected field $B_x$ is completely overwhelmed by the reconnecting field $B_y$. The gradient of the background outflow in the $y$ direction, whose length scale is $\\sim\\Lsheet$, is also unimportant because $k_{\\rm max}\\Lsheet\\gg1$ everywhere in the sheet. At the periphery of the sheet, for $y_0 > y_{0,\\rm crit}$, where $y_{0,\\rm crit}/\\Lsheet$ is equilibrium-dependent but otherwise $\\mathcal O(1)$, the current sheet becomes unstable to the Kelvin-Helmholtz (KH) instability driven by the velocity shear between the Alfv\\'enic reconnection outflow and the stationary upstream plasma. This occurs because the magnitude of the upstream magnetic field is a decreasing function of the outflow coordinate $y$ and eventually becomes smaller than the outflow speed (which is an increasing function of $y/\\Lsheet$). At, and beyond, the Alfv\\'en Mach point, where this happens, the magnetic field can no longer stabilize the current sheet against the KH instability.  We find that the KH instability of the sheet can either be resistive (i.e., induce reconnection at $x=0$), or ideal (no reconnection), with lower values of $k\\Lsheet$ corresponding to the former, and larger values to the latter. The fastest growing KH mode, $k_{\\rm max}\\Lsheet\\sim S^{1/2}$ (i.e., $k_{\\rm max}\\deltacs\\sim 1$), is an ideal, non-reconnecting mode. This is because reconnection cannot occur at the fast rates required by the fastest growing KH mode. A useful analogy can be made with the Taylor (forced reconnection) problem~\\cite{hahm_forced_1985}: since there are two shear layers, one on each side of the current sheet, the KH instability of the sheet is conceptually similar to a situation where perturbations at distant walls attempt to drive  reconnection at a rational surface. In the Taylor problem, it is also found that the perturbations at the walls do not drive reconnection in the initial stage. However, the same analogy suggests that as the ideal KH mode evolves into the nonlinear regime, it will cause the upstream magnetic field to pile-up in the current layer, eventually leading to its reconnection. This KH-driven reconnection that occurs at  $y_0 > y_{0,\\rm crit}$ will give rise to a plasmoid chain, just as the ``pure'' plasmoid instability that is found at $y_0 < y_{0,\\rm crit}$. Therefore, in practice, it may be difficult to distinguish between the two situations.  It is also worth noting that the basic KH instability mechanism that we have described here is completely general, i.e., it should apply to {\\it any} reconnecting current sheet, not just to those that can be described by the reduced MHD framework that we have adopted here: the only ingredient it requires is the existence of an Alfv\\'enic Mach point somewhere along the layer. This should be a generic feature of most reconnecting current sheets. Whether the layer is collisional or collisionless may affect the dynamics of the KH instability, but its existence is not dependent on the plasma collisionality. In this respect, our findings may be related to recent observations of the KH instability in collisionless simulations of guide-field reconnection~\\cite{fermo_2012}.  Finally, the effect of viscosity on the plasmoid instability has been addressed via numerical integration of the linearised set of equations. Our results are that in the limit $Pm = \\nu/\\eta \\gg 1$, the fastest growth rate and wave-number of the plasmoid instability scale as: \\be \\label{gmax_largePm_2} \\gmax\\sim S^{1/4}Pm^{-5/8}\\sim \\Lsheet^{1/4}~V_A^{1/4} \\eta^{3/8} \\nu^{-5/8}, \\ee \\be \\label{kmax_largePm_2} \\kmax\\sim S^{3/8}Pm^{-3/16}\\sim \\Lsheet^{3/8}~V_A^{3/8} \\eta^{-3/16} \\nu^{-3/16}. \\ee We have not performed a rigorous analytical calculation of the plasmoid instability in this limit, but we have been able to recover these scalings in a non-rigorous way from known results on the visco-tearing and resistive-kink modes~\\cite{porcelli_viscous_1987}, via the rescaling of the background magnetic shear length $a\\rightarrow \\deltacs\\sim \\Lsheet S^{-1/2}Pm^{1/4}$~\\cite{park_reconnection_1984}. Although these scalings are only expected to apply for $S\\gg \\Scrit,~Pm\\gg1$, where $\\Scrit$ is the critical value of the Lundquist number for the current sheet to be plasmoid-unstable, they lead to the prediction that \\be \\label{Scrit_largePm} \\Scrit\\sim 10^4 Pm^{1/2}, \\qquad Pm\\gg1. \\ee This result, as well as those of \\eqs{gmax_largePm_2}{kmax_largePm_2} are concrete predictions that can be tested in direct numerical simulations of MHD reconnection in the large magnetic Prandtl number regime."
    },
    "1208/1208.2963_arXiv.txt": {
        "abstract": "We present a theoretical framework in which bound stellar clusters arise naturally at the high-density end of the hierarchy of  the interstellar medium (ISM). Due to short free-fall times, these high-density regions achieve high local star formation efficiencies, enabling them to form bound clusters. Star-forming regions of lower density remain substructured and gas-rich, ending up unbound when the residual gas is expelled. Additionally, the tidal perturbation of star-forming regions by nearby, dense giant molecular clouds imposes a minimum density contrast required for the collapse to a bound cluster. The fraction of all star formation that occurs in bound stellar clusters (the cluster formation efficiency or CFE) follows by integration of these local clustering and survival properties over the full density spectrum of the ISM, and hence is set by galaxy-scale physics. We derive the CFE as a function of observable galaxy properties, and find that it increases with the gas surface density, from $\\Gamma\\sim1$\\% in low-density galaxies to a peak value of $\\Gamma\\sim70$\\% at densities of $\\Sigma_{\\rm g}\\sim10^{3}~\\msun~{\\rm pc}^{-2}$. This explains the observation that the CFE increases with the star formation rate density in nearby dwarf, spiral, and starburst galaxies. Indeed, comparing our model results with observed galaxies yields excellent agreement. The model is applied further by calculating the spatial variation of the CFE within single galaxies. We also consider the variation of the CFE with cosmic time and show that it increases with redshift, peaking in high-redshift, gas-rich disc galaxies. It is estimated that {up to} $30$--$35$\\% of all stars in the Universe once formed in bound stellar clusters. We discuss how our theory can be verified with {\\it Gaia} and ALMA, and provide possible implementations for theoretical work and for simulations of galaxy formation and evolution. ",
        "introduction": "\\label{sec:intro} The formation of stars from a turbulent interstellar medium (ISM) is caused by the fragmentation of hierarchically collapsing giant molecular clouds \\citep[e.g.][]{larson81,elmegreen04,maclow04,mckee07}. While this by necessity implies that stars generally do not form alone, it has been known for a long time that the star formation process does not exclusively produce stars in bound stellar clusters \\citep{elmegreen83,lada87}. Instead, some fraction of stars is born in unbound stellar associations over a broad range of physical scales \\citep[e.g.][]{blaauw64,clarke00,megeath04,portegieszwart10,gieles11,bastian12}, and in rare cases individual stars may even form in relative isolation \\citep{parker07,krumholz09,bressert12}. Surprisingly though, star clusters are often still considered to be a fundamental unit of star formation \\citep[e.g.][]{pflamm07,pfalzner09,assmann11}. Whereas galaxy formation studies augmented the scenario of monolithic collapse \\citep[see e.g.~the seminal paper by][]{eggen62} with the current picture of hierarchical galaxy formation \\citep[e.g.][]{white78,searle78,white91} several decades ago, the concept of star formation occurring in quantized, gravitationally bound systems has remained remarkably popular.  It was pointed out by \\citet{lada03} that in the Milky Way, embedded stellar groups constitute the vast majority of star formation and that these groups are about 10--20 times more numerous than gas-rid star clusters, suggesting that only a small fraction is capable of surviving the gas-embedded phase. With their review, the concept of {\\it infant mortality} became firmly established -- in this scenario, most (if not all) stars form in clusters, but the population of embedded clusters is then decimated due to the expulsion of residual gas by feedback \\citep[e.g.][]{tutukov78,hills80,lada84,adams00,geyer01,boily03,boily03b,bastian06b,goodwin06,baumgardt07,parmentier08}.  However, a recent study by \\citet[also see \\citealt{parker12b}]{bressert10} of the spatial distribution of young stellar objects (YSOs) in the solar neighbourhood has shown that star formation follows a continuous spectrum of number densities, indicating that there is no separate or critical density scale for star cluster formation imprinted in the star formation process. Their results imply that surface density thresholds used for the identification of stellar clusters \\citep[e.g.][]{lada03,jorgensen08,gutermuth09} are arbitrary, and do not correspond to a physical scale. It is thus impossible to conclude that the majority of stars form in clusters by dividing the density spectrum of star formation at a certain density. In other work, it has been shown that YSOs are strongly correlated with the hierarchically structured ISM \\citep{testi00,allen07,gutermuth11,bressert12b}, and that the impact of gas expulsion on bound stellar structure in a dynamical environment resulting from such a hierarchy is modest, if present at all \\citep{kruijssen12,girichidis12,cottaar12}. Rather than the scenario of star cluster formation in which infant mortality (i.e.~gas expulsion) plays an important role, these results favour a framework in which the observed scarcity of bound, gas-rid clusters is a simple result of the star formation process: only some fraction of star formation reaches the densities required to attain high star formation efficiencies and result in bound stellar clusters, while the rest forms in a more dispersed fashion throughout the natal cloud \\citep[e.g.][]{elmegreen01}.  The fraction of star formation occurring in bound stellar clusters is often quantified as the {\\it cluster formation efficiency} \\citep[CFE or $\\Gamma$, see e.g.][]{bastian08,goddard10,adamo11,silvavilla11}. The CFE is a crucial quantity in many respects. Not only can it enable a better understanding of the star formation process itself \\citep[e.g.][]{elmegreen02}, but it is also a key ingredient for work that aims to trace the (star) formation histories of galaxies using their star cluster populations \\citep[e.g.][]{larsen01,bastian05,smith07,konstantopoulos09,fedotov11}, or studies of star cluster disruption \\citep[e.g.][]{gieles05}. On galactic scales, the CFE may be used to infer whether the most massive stars inject their feedback energy within the same bubbles and hence how efficiently feedback couples to the ISM \\citep[cf.][]{strickland99,krause12}. Additionally, the CFE is an important parameter in numerical simulations on galactic or cosmological scales that aim to model the assembly history of star cluster systems \\citep[e.g.][]{prieto08,kruijssen11,kruijssen12c}. Thusfar, such simulations have had to assume that the CFE is constant throughout cosmic time.  Observational studies of extragalactic cluster populations at ages older than the embedded phase suggest that the CFE increases with the star formation rate surface density of the host galaxy \\citep{larsen00b,goddard10,adamo11,silvavilla11}, but at present there is no theoretical understanding of the variation of the CFE with the galactic environment. It is notoriously hard to model the formation of stellar clusters in numerical simulations due to the large dynamic range that needs to be covered to resolve the necessary physics. Current state-of-the art simulations can model systems up to masses of a few $10^3~\\msun$ \\citep[e.g.][]{tilley07,bonnell08,krumholz12}, but this is still insufficient by several orders of magnitude to address the formation of cluster populations on the desired galactic scales.  The aim of this paper is to derive and apply a theory of cluster formation that is based on simple analytical considerations. In part, it will follow the recent approach by \\citet{elmegreen08}, who showed that the CFE can be related to the density spectrum of the ISM, but did so by defining a critical density (or pressure) for cluster formation. While not suitable to do quantitative predictions due to an arbitrarily defined density threshold, the work is a big step forward from the idealized picture in which centrally concentrated gas is expelled from spherically symmetric, bound clusters. In the present paper, the avenue suggested by \\citet{elmegreen08} is explored further and expanded to provide a self-consistent framework for star cluster formation in galaxies.  The paper is organized as follows. In \\S\\ref{sec:clform} the required physical mechanisms are discussed that should be included in a theory of cluster formation. The model is derived and presented in \\S\\ref{sec:model}--\\ref{sec:cfe}. In \\S\\ref{sec:param} the parameter space is explored and it is addressed how the CFE varies among different galaxy types. The calculated CFEs are compared to observed cluster populations in \\S\\ref{sec:obs}, giving excellent agreement. In \\S\\ref{sec:spatial} the spatial variation of the CFE within single galaxies is addressed. In \\S\\ref{sec:hubble} the variation of the CFE with cosmic time is discussed, accounting for the evolution of galaxy properties in a cosmological context. A discussion of possible caveats is presented in \\S\\ref{sec:disc}, together with an outlook to the observational verification of the model, and potential applications in theoretical and numerical work. The conclusions of this paper are given in \\S\\ref{sec:concl}. We provide Fortran and IDL routines for calculating the CFE with the model of this paper at http://www.mpa-garching.mpg.de/cfe (see Appendix~\\ref{sec:appsupp}). ",
        "conclusions": "\\label{sec:concl} We have described and applied an analytic theory of the cluster formation efficiency (CFE), i.e.~the fraction of star formation occurring in bound stellar clusters. The conclusions of the paper are as follows. \\begin{enumerate} \\item The presented theory of cluster formation provides a framework in which the galaxy-scale environment is related to the properties of star-forming regions. Bound stellar clusters naturally arise from the high-density end of the density spectrum of the interstellar medium (ISM). Due to short free-fall times, these high-density regions can achieve high star formation efficiencies (SFEs). This makes them insensitive to gas expulsion and enables them to form bound clusters. In regions of lower density, the SFE is lower and the collapse into spherically symmetric structure is not completed. As a result, such regions (or associations) remain gas-rich and are unbound upon gas expulsion. (\\S\\ref{sec:ismpdf}--\\ref{sec:fbound}) \\item In the picture described here, gas expulsion does not affect centrally concentrated stellar clusters like in the classical scenario of `infant mortality', but inhibits the merging of hierarchically structured stellar groups and associations. (\\S\\ref{sec:fbound} and Appendix~\\ref{sec:appbound}) \\item Additionally, we have included the `cruel cradle effect', which represents the tidal perturbation of star-forming regions by surrounding GMCs. This is shown to give a second-order decrease of the CFE in normal disc galaxies, and is most important in high-density galaxies. (\\S\\ref{sec:cce}) \\item The CFE is obtained by integrating the probability distribution function (PDF) of the densities at which star formation proceeds, and the corresponding part thereof that is both bound locally and sufficiently dense to survive the cruel cradle effect. (\\S\\ref{sec:cfe} and Figure~\\ref{fig:xpdf}) \\item Our theoretical framework allows the prediction of the CFE as a function of galaxy-scale observables: the gas surface density $\\Sigma_{\\rm g}$, the angular velocity $\\Omega$, and the Toomre $Q$ parameter. For use in numerical simulations, a local formulation is also given, which uses the gas volume density $\\rho_{\\rm loc}$, the gas velocity dispersion $\\sigma_{\\rm loc}$, and the sounds speed $c_{\\rm s,loc}$. (\\S\\ref{sec:cfe} and \\S\\ref{sec:num}) \\item The model predicts that only some fraction (1--70\\%) of all star formation results in bound stellar clusters. This CFE increases with the gas surface density of the galaxy, because in denser galaxies a larger fraction of the ISM is pushed into the high-density regime where bound cluster formation occurs. However, at surface densities $\\Sigma_{\\rm g}\\ga10^{3}~\\msun~{\\rm pc}^{-2}$ the CFE is limited by the cruel cradle effect. For $\\Sigma_{\\rm g}\\ga10^{2.5}~\\msun~{\\rm pc}^{-2}$, the CFE decreases with increasing angular velocity $\\Omega$, while at lower densities it is generally insensitive to $\\Omega$. The CFE increases with Toomre $Q$, and hence with the stability of the galaxy disc. (\\S\\ref{sec:cfesflaw}, \\S\\ref{sec:cfeQ}, and \\S\\ref{sec:summparam}) \\item We find remarkable agreement between our theoretical model and the observed relation between the CFE and the star formation rate density $\\Sigma_{\\rm SFR}$ in dwarf, spiral and starburst galaxies. Also when modeling individual galaxies, the model predictions are in excellent agreement with the observed CFEs. (\\S\\ref{sec:obs}) \\item The model is applied to investigate the spatial variation of the CFE within galaxies, using parameter sets that resemble the Milky Way, M31, and M51. At the low densities of the Milky Way and M31, the CFE simply follows the relatively flat surface density profile of the gas. However, for the exponential gas density profile of M51, the CFE increases towards the centre of the galaxy, but (at high densities) does so less than the gas density profile. (\\S\\ref{sec:spatial}) \\item We use samples of nearby and high-redshift galaxy discs and starbursts to predict the evolution of the CFE across a Hubble time of star formation. For galaxy discs, the CFE is found to decrease from $\\Gamma=30$--$70$\\% to $\\Gamma=5$--$10$\\% between redshifts $z\\sim2$ and $z=0$, whereas starburst galaxies exhibit a moderate decrease of $\\Gamma=20$--$60$\\% to $\\Gamma=10$--$50$\\% over the same redshift range. High-redshift disc galaxies may thus have been the most efficient cluster factories in the history of the Universe, although presently the CFE is generally higher in starburst galaxies than in discs. We estimate that {up to} $\\Gamma_{\\rm univ}=30$--$35$\\% of all stars in the Universe was formed in bound stellar clusters. (\\S\\ref{sec:hubble}) \\item The predictions of our theoretical framework are well-suited for verification with {\\it Gaia} or ALMA, and also for application to observational studies of cluster populations. We provide specific recipes for use of the model in theoretical and numerical work, such as numerical simulations of galaxies and their star cluster populations, the feedback efficiency in cosmological simulations, and semi-analytic galaxy models. (\\S\\ref{sec:applobs} and \\S\\ref{sec:applth}) \\end{enumerate} We have presented a new model that describes how populations of bound stellar clusters form, and that predicts which fraction of all star formation is constituted by these clusters (the cluster formation efficiency or CFE). The model enables the calculation of the CFE over a range of galactic environments, and shows that star clusters are no fundamental unit of star formation, but instead are a possible outcome."
    },
    "1208/1208.0157_arXiv.txt": {
        "abstract": "Numerical simulations of the supernova (SN) neutrino self-induced flavor conversions, associated with the neutrino-neutrino interactions in the deepest stellar regions, have been typically carried out assuming the ``bulb-model''. In this approximation, neutrinos are taken to be emitted half-isotropically by a common neutrinosphere. In the recent Ref.~\\cite{Mirizzi:2011tu}  we have   removed this assumption by introducing flavor-dependent angular distributions for SN neutrinos, as suggested by core-collapse simulations. We have found that in this case  a novel multi-angle instability in the self-induced flavor transitions can arise. In this work we perform an extensive study of  this  effect, carrying out  a linearized  flavor stability analysis   for different SN neutrino energy fluxes and angular distributions, in both normal and inverted neutrino mass hierarchy. We confirm that spectra of different $\\nu$ species which cross in angular space (where $F_{\\nu_e}=F_{\\nu_x}$ and $F_{\\bar\\nu_e}=F_{\\bar\\nu_x}$) present a significant enhancement of the flavor instability, and a shift of the onset of the flavor conversions at smaller radii with respect to the case of an isotropic neutrino emission. We also illustrate how a qualitative (and sometimes quantitative) understanding of the dynamics of these systems follows from a stability analysis. ",
        "introduction": "Flavor conversions of neutrinos emitted  from  core-collapse supernovae (SNe) represent a diagnostic tool to get crucial information about their  mixing parameters and the SN dynamics in the deepest stellar regions~\\cite{Raffelt:2010zza}. In particular, SNe are a unique laboratory to probe neutrino oscillations in extreme conditions. Supernova  neutrinos can  interact not only with the stellar medium via the Mikheyev-Smirnov-Wolfenstein (MSW) effect~\\cite{Matt,Dighe:1999bi},  but also with other neutrinos  ($\\nu$) and  antineutrinos (${\\overline\\nu}$) as well. It was pointed out that large $\\nu$ densities in the deepest stellar regions can result in significant coherent $\\nu$--$\\nu$ forward scatterings~\\cite{Pantaleone:1992eq,Qian:1994wh}. A few years ago it was discovered that $\\nu$ self-interactions can give rise to collective $\\nu$ flavor oscillations inside the SN~\\cite{Duan:2006an,Hannestad:2006nj,Pehlivan:2011hp} (see~\\cite{Duan:2010bg} for a recent review). The most important observational consequence of this collective behavior is a swap of the $\\nu_e$ and ${\\overline\\nu}_e$ spectra with the non-electron $\\nu_x$ and ${\\overline \\nu}_x$ spectra in certain energy ranges~\\cite{Fogli:2007bk,Dasgupta:2009mg}. It has been argued that ``spectral splits'' at the edge of each swap interval would be observable in the high-statistics $\\nu$ signal from the next galactic SN, allowing to get crucial information about the unknown $\\nu$ mass ordering (see, e.g.,~\\cite{Duan:2007bt}).   Self-induced oscillation  effects crucially depend on the inner  boundary conditions fixed for the further flavor evolution. Indeed, these non-linear flavor conversions are associated to instabilities in the flavor space, that develop around the crossing points in the energy spectra of the different $\\nu$ species (where $F_{\\nu_e}=F_{\\nu_x}$ and $F_{{\\bar\\nu}_e}=F_{{\\bar\\nu}_x}$)~\\cite{Dasgupta:2009mg}. The number and the position of the crossing points  depend on the ordering of the original SN $\\nu$ fluxes. Since this latter can   change during the different post-bounce stages, significant temporal variations are expected in the pattern of the spectral splits~\\cite{Dasgupta:2009mg,Duan:2010bf,Mirizzi:2010uz}. An important layer of complication in this description is associated to the the current-current nature of the $\\nu$-$\\nu$  weak-interaction Hamiltonian. This implies that the interaction energy between neutrinos of momenta ${\\bf p}$ and ${\\bf q}$ is proportional to $(1 - {\\bf v_p} \\cdot {\\bf v_q})$, where ${\\bf v_p}$ is the neutrino velocity~\\cite{Pantaleone:1992eq,Qian:1994wh,Sawyer:2008zs}. In a non-isotropic medium this velocity-dependent term would not average to zero, producing a different refractive index for neutrinos propagating on different trajectories. This is the origin of the so-called ``multi-angle effects,'' which in some case can dramatically affect the development of the self-induced flavor conversions, producing a quick flavor decoherence~\\cite{Raffelt:2007yz,EstebanPretel:2007ec}  or suppressing flavor conversions otherwise possible for an   a isotropic neutrino emission~\\cite{Duan:2010bf,Mirizzi:2010uz}. Moreover, it has been shown that the presence of the ordinary matter background would cause a multi-angle suppression of the collective oscillations, when the matter density dominates over the neutrino one~\\cite{EstebanPretel:2008ni}. This situation is expected to occur during the early times accretion phase~\\cite{Chakraborty:2011nf,Chakraborty:2011gd,Sarikas:2011am,Saviano:2012yh,Sarikas:2012vb}.  The characterization of multi-angle effects is then a key-ingredient to obtain a realistic description of the self-induced neutrino flavor conversions. In this context, it is expected that the $\\nu$ angular distributions at emission would play an important role in determining the $\\nu$-$\\nu$ interaction strength. Until recently numerical simulations for the SN $\\nu$ flavor conversions  have been based on the so-called ``bulb model'' (see, e.g,~\\cite{Duan:2006an,Fogli:2007bk}), where $\\nu$'s of different species are considered as emitted ``half-isotropically'' (i.e. with  all outward-moving angular modes equally occupied and all the backward-moving modes empty) by a common spherical ``neutrinosphere,'' in analogy with a blackbody emission. However,  realistic supernova simulations show that $\\nu$  angular distributions at decoupling are far from being half-isotropic and, above all, are flavor-dependent (see, e.g.,~\\cite{Sarikas:2011am,Ott:2008jb,Dasgupta:2011jf}). The presence of non-trivial  angular distributions was claimed in~\\cite{Sawyer:2005jk}  to produce a novel multi-angle instability in the self-induced flavor evolution of a toy model of $\\nu$ gas. However, that conclusion was based on an analysis performed with a small number of angular modes, making challenging to rely on this claim to infer conclusions for  the realistic SN $\\nu$ case without a dedicated large-scale numerical study. Triggered also by this warning, in~\\cite{Mirizzi:2011tu} we performed numerical simulations of the self-induced flavor conversions with non-trival $\\nu$ angular distributions, finding remarkable effects on the flavor evolution. In particular, when flavor-dependent angular distributions  lead to crossing points in the angular spectra of different $\\nu$ species,  a new multi-angle instability can develop, in analogy to the known instability triggered by crossing points  in the energy domain. We find cases in which this  multi-angle instability can  shift the onset of the flavor conversions toward  low-radii and  produce a smearing of the splitting features found with trivial $\\nu$ emission models. In order to achieve a semi-analytical understanding of the phenomenon, in that Letter we proposed to apply to our problem the linearized stability analysis recently worked out in~\\cite{Banerjee:2011fj}. By seeking an exponentially growing solution in the eigenvalue equations associated with the linearized equations of motion for the neutrino ensemble, this method allowed us to determine the onset of the flavor conversions. In a specific scenario, we compared the growth of this mode in the half-isotropic case with different cases with non-trivial $\\nu$ angular distributions, finding a significant enhancement of the instability in these latter cases.  The purpose of this follow-up work is to take a closer look to the multi-angle instability, triggered by non-trivial angular distributions. In particular we aim to use the stability analysis to perform an investigation on the dependence of this effect on the initial SN $\\nu$ flux ordering, and on the $\\nu$ mass hierarchy. Here is the plan of our work. In Sec.~\\ref{setup} we  introduce the setup for the flavor-stability analysis, describing the non-linear equations for the $\\nu$ flavor  evolution in SNe, and the consistency equations coming from their linearization. In Sec.~\\ref{stability} we present our models for the supernova neutrino emission for cases with a different flux ordering and  angular distributions. For these different cases we show  the results of the stability analysis in the two neutrino mass hierarchies. Finally in Sec.~\\ref{conclusions} we comment on our results and we conclude. ",
        "conclusions": "\\label{conclusions}  Self-induced flavor conversions for SN neutrinos are associated to instabilities in the flavor space. Recently, a linearized stability analysis of the SN $\\nu$ equations of motion has been proposed as diagnostic tool to track the emergence of these  instabilities~\\cite{Banerjee:2011fj}. In our work we have applied this technique to study the multi-angle instability associated with non-trivial neutrino angular distributions, extending the results presented in~\\cite{Mirizzi:2011tu} to different SN $\\nu$ flux orderings and  $\\nu$ mass hierarchies. We confirm that enhanced instabilities can occur,  when the angular spectra present crossing points. Also, the onset of the flavor conversions typically shifts toward small radii, with respect to the half-isotropic case.  We cannot but emphasize that neutrino angular distributions should be taken into account in supernova neutrino phenomenological studies. For example,  as pointed out in~\\cite{Mirizzi:2011tu}, one can expect a smearing of the splitting features observed in the half-isotropic case, with resulting $\\nu$ fluxes showing less significant spectral differences. This tendency toward spectral equalization would challenge the detection of further oscillation signatures, like the ones associated with the Earth crossing of SN $\\nu$'s (see, e.g.,~\\cite{Lunardini:2001pb,Borriello:2012zc}). On the other hand, a sufficiently low-radii onset on the flavor conversions might imply an interesting impact on the $r$-process nucleosynthesis in SNe~\\cite{Duan:2010af}. Eventually, if $\\nu$ oscillations develop too close to the neutrinosphere, they might invalidate the $\\nu$ transport paradigm in SNe that ignores  $\\nu$ conversions.   It is clear that future, state-of-the-art predictions of the neutrino flavor evolution in SNe would require as accurate as possible input from hydrodynamical simulations, providing the flavor, angular, energy and time structure of the fluxes. Here we showed that a preliminary flavor stability analysis presents some key advantages over a brute force numerical integration. While such a perturbative analysis cannot be used to deduce the eventual fate of the neutrino ensemble in flavor space, it can be useful to identify the situations where angular distributions  play a role in the flavor evolution.  Once these cases are identified, it would be mandatory to perform large scale numerical simulations of the complete flavor evolution in order to determine the final oscillated neutrino spectra. With current level of sophistication in SN simulations, we believe that such studies will be mandatory to achieve a better characterization of the non-linear neutrino flavor evolution during a stellar gravitational collapse and a more realistic description of the neutrino signal from a future galactic SN."
    },
    "1208/1208.1960_arXiv.txt": {
        "abstract": "We present the results of the comparative analysis of the most known semi-empirical and empirical spectral atlases that was carried out using the data from the WBVR photometric catalogue. The results show that standard error of synthesized stellar magnitudes calculated with SEDs from best spectral atlases reaches 0.02 mag. It has been also found out that some of modern spectral atlases are burdened with significant systematic errors. The agreement for the 5000-10000 A spectral range is rather satisfactory, while there are problems for wavelengths shorter than 4400 A. ",
        "introduction": "\\label{s:intro}  For a great variety of astrophysical applications it is extremely important to know (at least relative) spectral energy distribution (SED) for \\begin{itemize} \\item as many stars as possible, and \\item as many types of stars as possible. \\end{itemize} Semi-empirical spectrophotometric atlases are designed to meet the latter requirement, meanwhile to satisfy the former one a large number of empirical atlases are constructed.  The Asiago Database of Spectroscopic Databases (Sordo \\& Munari, 2006) includes about 300 empirical stellar spectral atlases, published from the late sixties to 2005. The majority of them covers a very restricted spectral range or/and contains too small number of stars. Less than a dozen of atlases contain large enough number of objects (at least hundreds) and provide SEDs for spectral range of at least 3000-9000 A. Among them are Pulkovo Spectrophotometric Catalogue, Alma-Ata atlas, NGSL, ISO-SWS atlas and some others. However, a comparison of data from these atlases for common stars shows numerous and significant disagreements, especially for UV spectral range.  On the other hand, multicolor high-precision photometry for much larger number of stars is available in different photometric systems. The most precise photometry can be found in Hipparcos/Tycho-2 catalogues and in ground-based WBVR catalogue of northern stars. So, the reliability of data in stellar spectral atlases can be checked by a comparison of magnitudes, calculated with methods of synthetic photometry, with catalogued data.  Here we present the results obtained using the observational data from the Catalogue of WBVR magnitudes of Bright Northern Stars (Kornilov et al. 1991). This catalogue provides the four-colour WBVR photoelectric magnitudes for 13586 northern sky objects ($\\delta>-15$) obtained at the high-altitude observatory in Kazakhstan using a four-channel photometer attached to the 0.5m reflector. The mean error in V ($\\sigma$V) for nonvariable stars in the catalogue is $0.003^m$. Limiting stellar magnitude is $V=7.2^m$. The catalogue is complete up to $V=7^m$. The WBVR catalogue can be considered as a photometric catalogue of Hipparcos level accuracy. ",
        "conclusions": "  \\begin{itemize} \\item Standard error of synthesized stellar magnitudes, calculated with SEDs from best spectral atlases, reaches 0.02 mag. \\item Some of modern spectral atlases are burdened with significant systematic errors. \\item SEDs from majority of atlases show satisfactory agreement for the 5000-10000 A spectral range, but problems for wavelength shorter than 4400 A remain. \\end{itemize}"
    },
    "1208/1208.6547_arXiv.txt": {
        "abstract": "{The origin of large-scale magnetic fields in cosmic structures and the intergalactic medium is still poorly understood. We explore the effects of non-minimal couplings of electromagnetism on the cosmological evolution of currents and magnetic fields. In this context, we revisit the mildly non-linear plasma dynamics around recombination that are known to generate weak magnetic fields. We use the covariant approach to obtain a fully general and non-linear evolution equation for the plasma currents and derive a generalised Ohm law valid on large scales as well as in the presence of non-minimal couplings to cosmological (pseudo-)scalar fields. Due to the sizeable conductivity of the plasma and the stringent observational bounds on such couplings, we conclude that modifications of the standard (adiabatic) evolution of magnetic fields are severely limited in these scenarios. Even at scales well beyond a Mpc, any departure from flux freezing behaviour is inhibited.} ",
        "introduction": "\\label{sec:intro}  Our Universe is apparently magnetised on virtually all different length-scales probed by astronomical observations. The strength of magnetic fields in galaxies and galaxy clusters is of the order of $\\mu\\Gauss$~\\cite{Kronberg:1993vk,Han:2002ns,Govoni:2004as,Clarke:2000bz}, and a lower bound of the order of $10^{-15}\\, \\Gauss$ on coherent magnetic fields in the intergalactic medium has also been reported~\\cite{Tavecchio:2010mk,Ando:2010rb,Neronov:1900zz,Dolag:2010ni,Essey:2010nd}. The origin of these large-scale magnetic fields still remains unclear, particularly at coherence lengths stretching beyond a Mpc, which is the subject of the present work\\footnote{In this context, theoretical models operating during QCD or Electroweak phase transitions (recent representative works are~\\cite{Saveliev:2012ea,Tevzadze:2012kk,Jedamzik:2010cy,Caprini:2009pr,Urban:2009sw}) will not be useful as the magnetic power spectrum they produce is cut off by causality at much smaller scales.}. Speculative mechanisms, active at all stages of the evolution of the Universe, have been proposed which often require the breaking of the conformal invariance of electromagnetism (EM). For recent reviews see refs.~\\cite{Subramanian:2009fu,Kandus:2010nw,Widrow:2011hs,Ryu:2011hu}.  Besides these exotic possibilities, it is known that weak magnetic fields are generated around photon decoupling. In the pre-recombination era, the vorticities of electrons and protons evolve slightly differently yielding a net circular current that generates magnetic fields, an effect known as the Harrison mechanism~\\cite{Harrison:1970,Harrison:1973zz}. However, no magnetic fields are produced at first order in cosmological perturbations, even in the presence of active sources of vector perturbations~\\cite{Hollenstein:2007kg}. At second order in both perturbations and the tight coupling approximation of the photon-baryon interactions, the mechanism does generate magnetic fields of about $10^{-29}\\,\\Gauss$ on Mpc scales~\\cite{Berezhiani:2003ik,Matarrese:2004kq,Gopal:2004ut,Ichiki:2007hu,Kobayashi:2007wd,Takahashi:2007ds,Maeda:2008dv,Fenu:2010kh,Maeda:2011uq,Giovannini:2011tj}.  This result provides a motivation to take a closer look at the evolution of EM fields during the late Universe. Clearly, in standard Maxwell theory the high conductivity of the cosmic plasma causes the electric fields to be screened rapidly and the magnetic fields to freeze into the plasma and evolve adiabatically, i.e.\\ $B\\propto a^{-2}$. However, one can expect a different behaviour in case Maxwell theory is modified on large scales. Specific scenarios have already been explored \\cite{Lee:2001hj, Urban:2009sw, BeltranJimenez:2010uh}, but no systematic investigation has been carried out so far.  In this paper, we consistently examine whether the seed fields from the recombination era could be boosted at late times due to non-minimal (effective) interactions of EM. In particular, we allow for different extensions of EM to embrace couplings with a scalar or pseudo-scalar field, for instance, dark energy (DE) or a cosmological axion. In this context, we review the hydrodynamical evolution of the primordial plasma from electron-positron annihilation until today by carefully taking into account the effects of a possible non-minimal theory. To this end we employ the 1+3 covariant approach to cosmology that is fully non-linear, not restricted to a Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) background, and captures all relativistic effects. Only after linearising, for illustration, it will be convenient to switch to a description using 3-vectors w.r.t.\\ a comoving basis in a FLRW background.  First of all, in section~\\ref{sec:ohm} we provide a brief review of the plasma processes occurring from the pre-recombination era to today by looking at the evolution of currents, relevant on cosmological scales at second order. Then, we derive a generalised Ohm's law that is also valid in the context of non-minimally coupled EM. By inspecting the different timescales involved in the current evolution we briefly discuss its validity on very large (sub- and super-horizon) cosmological scales. This section should serve as a reference for cosmological plasma conductivity formulae and physics. In section~\\ref{sec:modified}, we introduce the generic form of the Lagrangian allowing for different non-minimal couplings and discuss their effects on the magnetic field evolution at large scales. Finally, in section~\\ref{TheEnd}, we summarise our results and discuss the conclusions. Appendix~\\ref{app:conv} lists our conventions, and appendix~\\ref{app:covariant} contains the details of the 1+3 covariant approach for the description of the fluid and current evolution in the plasma. ",
        "conclusions": "\\label{TheEnd}   Large-scale magnetic fields are present not only in bound cosmic structures such as galaxies and clusters, but also they appear to fill a large fraction of the intergalactic space. This fact is difficult to explain by means of standard astrophysical processes. Also the early Universe mechanisms of magnetogenesis are mostly not viable or not efficient, even if one allows for non-minimal EM. On the other hand, the non-linear dynamics of the cosmic plasma in the recombination era do source large-scale magnetic fields, even though with insufficient amplitudes. Taking this mechanism as a promising starting point, we have carefully revisited the non-linear dynamics in a more general setting, allowing for non-minimal couplings of EM to (effective) scalar fields, in the hope of boosting the magnetic seeds.  We have derived the fully general and non-linear evolution equation of the current, and specialised it to the photon-electron-proton plasma, relevant in the recombination era, taking into account all possible effects up to second order in perturbations. We paid particular attention to allowing for possible effects of non-minimal couplings of EM, refraining from applying the Maxwell and Einstein field equations that we want to modify. Subsequently, we have discussed the relevance of the different effects on the current and derived a cosmological Ohm's law that is valid even in non-minimal EM. We also point out that Ohm's law is valid, and the resistivity and conductivity are well defined, on large cosmological scales, even beyond the Hubble horizon, as long as super-horizon correlations were initially imprinted into the cosmic plasma (through inflation).  We have investigated several different types of couplings of EM to scalar and pseudo-scalar fields. Candidate fields which could do the job are for example the DE field, or a cosmological axion background. We have explored the possibilities of a running fine-structure constant, $f^2(\\phi)F^2$, a running axial coupling, $g(\\phi)F\\tilde{F}$, and an effective photon mass $m^2(\\phi)A^2$. In all three cases we conclude that current observational and experimental limits do not allow for additional amplification of EM fields on cosmological scales. Even at scales well beyond a Mpc the conductivity of the plasma is so high that it inhibits any departure from the flux freezing behaviour, resulting in the strength of the magnetic field to be about $3\\times10^{-29}\\,\\Gauss$ today."
    },
    "1208/1208.4297_arXiv.txt": {
        "abstract": "We give an overview of recent results for the nuclear equation of state and properties of neutron stars based on microscopic two- and three-nucleon interactions derived within chiral effective field theory (EFT). It is demonstrated that the application of Renormalization Group (RG) transformations allows efficient, controlled and simplified calculations of nuclear matter. ",
        "introduction": "Establishing an interparticle Hamiltonian, which is the fundamental ingredient for microscopic many-body calculations of atomic nuclei and nucleonic matter, is a difficult and on-going challenge for low-energy nuclear physics. Chiral EFT and RG methods make feasible a controlled description of nuclear interactions, grounded in QCD symmetries, which in turn makes possible the description of nuclei across the nuclear many-body landscape~\\cite{Epelbaum_RevModPhys, Bogner_review}. The two-body sector has been solved in the sense that various nucleon-nucleon interactions are available that reproduce low-energy scattering phase shifts to high accuracy. The unsettled frontier is three- and higher-body forces, which play a key role in many-nucleon systems.  Nuclear structure calculations are complicated due to the coupling of low to high momenta by nuclear interactions. A fruitful path to decoupling high-momentum from low-momentum physics is the Similarity Renormalization Group (SRG), which is based on a continuous sequence of unitary transformations that suppress off-diagonal matrix elements. The decoupling of momenta via the SRG has been demonstrated for nuclear nucleon-nucleon (NN) interactions~\\cite{Bogner_SRG_evolution} and very recently for the first time also for three-nucleon forces (3NF)~\\cite{Hebeler_3NF_evolution}. The decoupling in 3NF is illustrated in Fig.~\\ref{fig:3NF_decoupling}: at large resolution scales $\\lambda = \\infty$ the potential contains significant off-diagonal couplings. As we evolve to lower resolution, these couplings get successively suppressed and finally at $\\lambda = 1.5\\:\\rm{fm}^{-1}$ the non-perturbative features are substantially softened as we find only non-vanishing strength at small momenta and around the diagonal.  SRG-evolved potentials are automatically energy independent and the same transformations renormalize  all operators, including many-body operators, and the class of transformations can be tailored for effectiveness in particular problems. When evolving nuclear interactions to lower resolution, it is inevitable that many-body interactions and operators are induced even if initially absent. This might be considered as unnatural if nuclei could be accurately calculated based on only NN interactions, as was assumed for much of the history of nuclear structure calculations. However, chiral EFT reveals the natural scale and hierarchy of many-body forces, which dictates their inclusion in modern calculations of nuclei and nucleonic matter.  \\begin{figure}[t] \\includegraphics[width=0.95\\textwidth,clip]{429_Hebeler-f1.eps} \\caption{(from Ref.~\\cite{Hebeler_3NF_evolution}) Contour plot of the RG-evolved 3N potential as a function of the hypermomenta $\\xi$ and $\\xi'$ for a fixed hyperangle. Evolution proceeds from left to right, with softening evidence by the suppression of off-diagonal elements. } \\label{fig:3NF_decoupling} \\end{figure} ",
        "conclusions": ""
    },
    "1208/1208.6292_arXiv.txt": {
        "abstract": "We present a broad-band timing analysis of the accreting white dwarf system MV Lyrae based on data obtained with the \\Kepler\\ satellite. The observations span 633 days at a cadence of 58.8 seconds and allow us to probe 4 orders of magnitude in temporal frequency. The modelling of the observed broad-band noise components is based on the superposition of multiple Lorentzian components, similar to the empirical modelling adopted for X-ray binary systems. We also present the detection of a frequency varying Lorentzian component in the lightcurve of MV Lyrae, where the Lorentzian characteristic frequency is inversely correlated with the mean source flux. Because in the literature similar broad-band noise components have been associated to either the viscous or dynamical timescale for different source types (accreting black holes or neutron stars), we here systematically explore both scenarios and place constraints on the accretion disk structure. In the viscous case we employ the fluctuating accretion disk model to infer parameters for the viscosity and disk scale height, and infer uncomfortably high parameters to be accommodated by the standard thin disk, whilst in the dynamical case we infer a large accretion disk truncation radius of $\\approx10R_{WD}$. More importantly however, the phenomenological properties between the broad-band variability observed here and in X-ray binaries and Active Galactic Nuclei are very similar, potentially suggesting a common origin for the broad-band variability. ",
        "introduction": "Compact interacting binaries (CBs) are close binary systems usually consisting of a late-type star that transfers material onto a black hole (BH), a neutron star (NS) or a white dwarf (WD) via Roche-lobe overflow. With an orbital period on the order of hours, the donor star transfers material through the L1 point, which forms an accretion disk surrounding the compact object. The dynamics and physics governing the flow of matter accreting onto the compact objects is, however, still debated. The accretion disks in BH and NS binary systems (X-ray binaries or XRBs) emit most of their radiation in X-rays, whilst the accretion disks in accreting white dwarfs (cataclysmic variables or CVs) emit mostly in the optical/UV wavebands. This is a consequence of the gravitational potential well created by the central compact object: for CVs the inner-most edge of the accretion disk sits at a few thousand gravitational radii, whilst for galactic BH and NS the inner disk reaches down to a few gravitational radii.  Most CVs and XRBs are highly variable sources in X-rays and/or optical/UV. This variability has been associated with the accretion disk from the characteristic frequencies observed in the power spectral density (PSD). For example, the presence of periodic modulations slightly longer than the orbital period can be associated with positive superhumps caused by a tidal deformation in the accretion disk (\\citealt{whitehurst,osaki,lubow}). This has been modelled and observed in CVs (\\citealt{montgomery,wood}), as well as in XRBs (\\citealt{dono}).  Another example of similar timing characteristics between CVs and XRBs is the presence of both dwarf nova oscillations (DNOs) and quasi-periodic oscillations (QPOs). In both cases DNOs/QPOs appear in the PSD at a few tens of mHz (\\citealt{warner_qpo,pretorius_qpo}), whilst for XRBs they appear at few hectoHz, and are referred to as lower kHz oscillations (\\citealt{belloni}). The phenomenological similarity between the QPOs and DNOs observed in CVs and XRBs was first noted by \\cite{mauche} and \\cite{warner_qpo}, where the ratio of periods is $P_{QPO}/P_{DNO} \\approx 15$, and holds over 6 orders of magnitude in temporal frequency. The physical reason for this relation is not fully understood, although some suggestions have been proposed (\\citealt{belloni}), all involving accretion disk dynamics, but none seem to be able to explain all observations consistently (\\citealt{warner_qpo,pretorius_qpo}).  In addition to periodic or quasi-periodic signals, both CVs and XRBs possess an intrinsic aperiodic broad-band noise continuum, generally described as flickering in CVs. This red noise component produces a continuum that rises towards low frequencies. In the XRB context, this is usually modelled a by simple power-law (e.g. $1/f$) or by a sum of broad Lorentzian components (\\citealt{belloni}), with a characteristic break at $\\approx 10^{-3}$ Hz for CVs and $\\approx 1$ Hz for XRBs (\\citealt{revn1,belloni}). This component has been shown to display a linear rms-flux relation in XRBs and Active Galactic Nuclei (AGN, \\citealt{uttley1}), where the root-mean-square (rms) variability linearly scales with flux over a wide range of timescales. More recently, the same relation has been also observed in the CV system MV Lyrae, albeit at lower frequencies (\\citealt{scaringi}, hereafter Paper~I). The detection of the rms-flux relation in both CVs and XRBs strongly suggests that the broad-band variability originates within the accretion disk. Under certain assumptions about the origin of the variability and the physics of the mass-transfer process, the characteristics of the broad-band noise can be used to constrain the strength of the effective viscosity that drives the mass transfer through the disk (\\citealt{ID1,ID2}).  Motivated by the similarities between the periodic and aperiodic signals in CVs and XRBs, we here study the broad-band temporal frequency spectrum of the CV MV Lyrae in search of further similarities between the two classes of objects. MV Lyrae is one of 14 known CVs in the \\Kepler\\ field-of-view (FOV), and is classified as being a VY Scl novalike system, spending most of its time in a high state ($V\\approx12-13$), but occasionally (every few years) undergoing short-duration (weeks to months) drops in brightness ($V\\approx16-18$, \\citealt{hoard}). The reason for these sudden drops in luminosity is not clear, but \\cite{LP} have suggested star spots from the donor covering the L1 point inhibiting mass transfer. It is known however that MV Lyrae has an extremely low mass transfer rate at its minimum brightness ($3\\times10^{-13}M_{\\odot}/yr$, \\citealt{linnell,hoard}), where the WD is detected at $V\\approx18$ and dominates the emitted light. Furthermore, the orbital period of $3.19$ hours has been determined for the system, as well as a low inclination of $i\\sim11^o-13^o$ (\\citealt{SPT95}).  In Paper~I we analysed the high frequency (tens of minutes) broad-band variability for MV Lyrae with data obtained with the \\Kepler\\ satellite. In this paper we will analyse the same lightcurve at lower temporal frequencies, with particular emphasis on the varying broad-band components during the observation. MV Lyrae has been observed with \\Kepler\\ during a low-to-high luminosity transition, when its optical emission originates mostly from its nearly face-on disk.  In section \\ref{sec:data}, we briefly describe the \\Kepler\\ data acquisition, and the procedure we use to construct broad-band power spectral densities (PSDs). We fit the PSDs with a combination of Lorentzian shaped functions to characterise the observed broad-band noise components. Section \\ref{sec:results} shows our results, with particular emphasis on one frequency-varying broad-band noise component, which appears to correlate with the mean source flux. Finally, in section \\ref{sec:discussion}, we consider whether the inferred broad-band noise components can be associated with fluctuations on either the dynamical or viscous time-scales in the accretion disk. As we shall see, both of these interpretations lead to uncomfortable implications for the accretion disk structure. ",
        "conclusions": "We have presented an analysis of the broad-band frequency behaviour of the accreting WD MV Lyrae based on data obtained by the \\Kepler\\ satellite. We have shown how the complex PSD can be decomposed with a number of Lorentzian-shaped functions. We further searched for possible correlations between the characteristic frequencies, and found the first frequency varying QPO in a CV, where frequency is inversely proportional to mean source flux.  The characteristic frequencies associated with the fitted Lorentzians were used to explore the origin of variable emission in terms of viscous or dynamical processes as the two limiting cases. In the former case we infer extremely high values of both disk scale height $H/R\\ge0.3$, and viscosity $\\alpha\\ge0.1$, suggesting the existence a geometrically-thick disk. This result is potentially in line with the work of \\cite{knigge00,knigge04}, which have suggested the presence of a self-occulting accretion disk close to the WD. In the dynamical case we instead infer a large disk truncation radius of $\\approx10R_{WD}$, but the presence of other components undermines a dynamical interpretation, at least for the lowest frequency components. More importantly, the presence of the rms-flux relation observed in MV Lyrae (\\citealt{scaringi}) and other XRBs/AGN (\\citealt{uttley1}), potentially rule out dynamical effects as the source of variability (since the fast dynamical variability will be damped), but favours a viscous origin for the observed broad-band noise components, at least at the highest observed frequencies.  In summary both viscous and dynamical (and consequently thermal) timescales struggle to consistently explain the observed broad-band variability in MV Lyrae. However, our analysis seems to suggest, at least phenomenologically, that the mechanisms which give rise to the observed rms-flux relation(s) and characteristic frequencies observed here and in other XRBs/AGN need to occur in all accretion discs (whether they are thin or thick), possibly suggesting a similar physical origin for the variability in both types of systems.  The broad-band variability properties of MV Lyrae remain yet to be fully understood. More generally, the variability properties of both CVs and XRBs also remain an enigmatic observational feature. Although both types of systems possess an accretion disk which is in many ways similar, there has not yet been enough observational data to provide a complete comparative study of the broad-band variability properties. However, it is clear that the phenomenological properties between the broad-band variability observed in MV Lyrae and in X-ray binaries and Active Galactic Nuclei are very similar (when appropriately scaled), and are potentially driven by a common accretion mechanism."
    },
    "1208/1208.4404_arXiv.txt": {
        "abstract": "Understanding the dynamic solar chromosphere is fundamental in solar physics. Spicules are an important feature of the chromosphere, connecting the photosphere to the corona, potentially mediating the transfer of energy and mass. The aim of this work is to study the properties of spicules over different regions of the sun. Our goal is to investigate if there is more than one type of spicules, and how spicules behave in the quiet sun, coronal holes, and active regions. We make use of high-cadence and high-spatial resolution \\caii\\ H observations taken by \\emph{Hinode}/SOT. Making use of a semi-automated detection algorithm, we self-consistently track and measure the properties of 519 spicules over different regions. We find clear evidence of two types of spicules. Type I spicules show a rise and fall and have typical lifetimes of $150-400\\;$s and maximum ascending velocities of $15-40\\;\\kms$, while type II spicules have shorter lifetimes of $50-150\\;$s, faster velocities of $30-110\\;\\kms$, and are not seen to fall down, but rather fade at around their maximum length. Type II spicules are the most common, seen in quiet sun and coronal holes. Type I spicules are seen mostly in active regions. There are regional differences between quiet sun and coronal hole spicules, likely attributable to the different field configurations. The properties of type II spicules are consistent with published results of Rapid Blueshifted Events (RBEs), supporting the hypothesis that RBEs are their disk counterparts. For type I spicules we find the relations between their properties to be consistent with a magnetoacoustic shock wave driver, and with dynamic fibrils as their disk counterpart. The driver of type II spicules remains unclear from limb observations. ",
        "introduction": "Spicules can be seen almost everywhere at the solar limb. They are highly-dynamic, thin, jet-like features in the solar chromosphere. Spicules have been observed for a very long time and have been the subject of numerous reviews \\citep[\\emph{e.g.}][]{Beckers:1968,Beckers:1972,Suematsu:1995,Sterling:2000}. A fundamental part of the chromosphere, their origin and connection to the photosphere and corona have long been debated.  Perhaps the most interesting aspect about spicules is their potential to mediate the transfer of energy and mass from the photosphere to the corona. This potential has been recognized early \\citep{Beckers:1968, PneumanKopp:1978, Athay:1982}, but the lack of high-quality observations prevented a better understanding of spicules and their link to the corona \\citep{Withbroe:1983}. The advent of the \\emph{Hinode} Solar Optical Telescope \\citep[SOT,][]{Kosugi:2007,Tsuneta:2008,Suematsu:2008} provided a quantum leap in the understanding of spicules and their properties. With its seeing-free high spatial and temporal resolution observations of the chromosphere, \\emph{Hinode} showed much more dynamic spicules than previously thought and re-ignited the discussion on their potential to heat the corona \\citep{DePontieu:2007, DePontieu:2009, DePontieu:2011}. %  Crucial for the understanding of the origin and effects of spicules is a characterization of their properties. While there are several recent studies of the properties of spicules using \\emph{Hinode}/SOT \\citep{DePontieu:2007,Anan:2010,Zhang:2012}, they are either brief first results papers and/or cover only one or two data sets. A comprehensive study of the properties of spicules in different solar regions is missing, and that is the aim of this work. Our goal was to obtain statistically-significant measurements of the properties of spicules in three key regions: quiet sun, coronal holes and active regions. In this we make use of the wealth of high-quality data taken with \\emph{Hinode}/SOT, and of automated detection algorithms.  The outline of this paper is as follows. A description of the observations is made in Section~\\ref{sec:obs}, followed by a description of the data analysis and automated spicule detection in Section~\\ref{sec:spic_det}. Our results are presented in Section~\\ref{sec:results}, and in Section~\\ref{sec:discussion} we discuss their consequences, including the types of spicules found and in which regions, and a comparison with other results from the literature. Finally, our conclusions are summarized in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Measuring limb spicules is a daunting task. Providing long time series of seeing-free data, \\emph{Hinode} is arguably the best observatory for characterizing spicules. Using a semi-automated procedure we obtained a consistent and statistically-significant set of spicule measurements over quiet sun, coronal holes, and active regions.  We find strong evidence of two populations of spicules, type I and type II, in agreement with \\citet{DePontieu:2007} and in contradiction of \\citet{Zhang:2012}. The typical properties of type I spicules are lifetimes of $150-400\\;$s, maximum ascending velocities of $15-40\\;\\kms$, maximum heights of $4-8\\;$Mm, inclinations of $0-40\\;$degrees, scale heights of $1.5-2.5\\;$Mm, and maximum transverse speeds of $5-30\\;\\kms$. Type II spicules, on the other hand, have much shorter lifetimes of $50-150\\;$s, larger ascending velocities of $30-110\\;\\kms$, maximum heights of $3-9\\;$Mm (longer in coronal holes), inclinations of $0-30\\;$degrees, scale heights of $1.8-3\\;$Mm in quiet sun and $2.4-3.6\\;$Mm in coronal holes, and slightly larger maximum transverse speeds of $4-40\\;\\kms$. Type I spicules rise and fall with a parabolic trajectory, while type II spicules rise linearly until their maximum length, and then dissipate quickly, typically over their whole length and suggesting that they are heated out of the \\caii\\ passband. We find that type II spicules are dominant in the quiet sun and coronal holes, with type I spicules being found mostly in active regions.   Our results differ from the classical views on spicules by \\citet{Beckers:1968, Beckers:1972}, who reported spicules with long lifetimes and slow velocities -- resembling the properties that we measure for type I spicules. However, these comprehensive reviews include results from quiet sun and coronal hole observations that most likely included both type I and type II spicules. This puzzling discrepancy between our results and classical spicules is addressed in \\citet{Pereira:2012spiclett}, where it is found that the lower resolution of earlier studies can affect the spicule measurements, and by degrading \\emph{Hinode} data to similar conditions of these studies, we measure spicule properties that agree with classical spicules. Thus, we believe that this difference stems from the higher quality observations available today.    The observations of type I spicules are consistent with a magnetoacoustic shock wave driver \\citep{DePontieu:2004,Hansteen:2006,Heggland:2007}, and that dynamic fibrils and mottles are their disk counterparts \\citep{DePontieu:2007DF,vdVoort:2007}. Given the relations between the measured decelerations, velocities, and lifetimes, our results support this view.  The ubiquity of type II spicules lends credence to the hypothesis that they can supply the corona with energy and mass, even if only a small percentage of them reach coronal heights \\citep{PneumanKopp:1978,DePontieu:2009,DePontieu:2011}. Furthermore, the suggestion by \\citet{Langangen:2008} that RBEs are the disk-counterpart of type II spicules has been underpinned by \\citet{vdVoort:2009} and \\citet{Sekse:2012}. Nevertheless, more detailed observations on the origins of type II spicules and improved numerical models are necessary for one to grasp their complex interaction from the photosphere to the corona."
    },
    "1208/1208.5222_arXiv.txt": {
        "abstract": "External errors of effective temperatures of stars for selected libraries are estimated from data intercomparisons. It is found that the obtained errors are mainly in a good correspondence with the published data. The results may be used to homogenize the effective temperatures by averaging the data (with the weights inversely proportional to the squared errors) from independent sources. ",
        "introduction": "Numerous spectral and photometric stellar catalogues of the basic atmospheric parameters ($T_{\\rm eff}$, $\\log g$, [Fe/H]) are widely used for decoding the structure, evolutionary stage and chemical enrichment history of the Galaxy. The rapidly growing number of such catalogues has imposed a need for refining procedures of merging these  stellar data into a single  homogenized catalogue. In our recent paper \\citep{2011BaltA..20...91M}, the technique has been applied where some homogeneous samples from selected catalogues of the $T_{\\rm eff}$ values were treated by combining them in triples, quadruples, quintuples and pairs for the stars in common to determine their external errors from data intercomparisons. The $T_{\\rm eff}$ values are then averaged (with the weights inversely proportional to the squared errors) to produce an extended mean homogenized catalogue. A somewhat different procedure of homogenization was used by \\citet{2007MNRAS.374..664C} where only pairs of stars were treated and heterogeneity of the published errors (or other quality data) of the parameters inside every used sample was not taken into account.  The most important and detailed information for stars is available in stellar spectral libraries with medium to high resolutions and good coverages of the Herzsprung-Russell diagram and metallicity range. Some libraries also contain the results of an efficient parametrization of $T_{\\rm eff}$, $\\log g$, [Fe/H]. Few libraries are considered here whose published $T_{\\rm eff}$ are compared with the  $T_{\\rm eff}$ from an independent extensive homogeneous sample with the external ${\\sigma T_{\\rm eff}}$ taken from our previous analysis \\citep{2011BaltA..20...91M}. These comparisons allow to estimate the ${\\sigma T_{\\rm eff}}$ for the considered spectral libraries, and to use the results for homogenizing the $T_{\\rm eff}$ values. ",
        "conclusions": "The technique of estimating the errors of catalogues from data intercomparisons has been applied to some samples of the published $T_{\\rm eff}$ for spectral libraries. The results will be used for producing an extended mean homogenized catalogue of  $T_{\\rm eff}$. The same approach may be applied also for the treatment of other kinds of data. The stars with the most reliable parameters  may then serve as templates in classification. Support from the Estonian Science Foundation (grant No. 7765) is acknowledged."
    },
    "1208/1208.0294_arXiv.txt": {
        "abstract": "The Kerr-Newman black hole solution can be constructed straightforwardly as the unique solution to the boundary value problem of the Einstein-Maxwell equations corresponding to an asymptotically flat, stationary and axisymmetric electro-vacuum spacetime surrounding a connected Killing horizon. ",
        "introduction": "It was shown in Ref.~\\refcite{m12} that the complex Ernst potentials\\cite{ernst68} on the symmetry axis of an asymptotically flat, stationary and axisymmetric electro-vacuum spacetime surrounding a connected Killing horizon (including the case of a degenerate horizon) can be constructed uniquely by means of the inverse scattering method\\cite{bv01, rfe}. This provides a new proof of the Kerr-Newman black hole uniqueness since the full solution can straightforwardly be derived from the axis data, cf.~Ref.~\\refcite{he81}. Here we present a particular method for obtaining the Ernst potentials off the axis, see Sec.~4. ",
        "conclusions": ""
    },
    "1208/1208.2772_arXiv.txt": {
        "abstract": "We present our near-infrared (NIR) imaging observations of the two  neutron star low mass X-ray binaries XMMU J174716.1-281048 and SAX J1806.5-2215 obtained using the PANIC instrument on the 6.5-meter Magellan telescope and the WHIRC instrument on 3.5-meter WIYN telescope. Both sources are members of the group of faint to very-faint X-ray binaries and  are undergoing very long X-ray outbursts  since 2003 and 2011, respectively-   `the quasi-persistent X-ray binaries'. The goal of our  observations is to identify the NIR counterpart of both sources. We identified two NIR stars  consistent with the  Chandra X-ray error circle of XMMU J174716.1-281048 and one for SAX J1806.5-2215. We studied the magnitude variations of the possible counterparts with respect to the UKIRT NIR Galactic plane observations. For  XMMU J174716.1-281048, we also  investigated the candidate counterparts using the color-color diagram and spectral energy distribution. We observed large variability in one NIR star having position consistent with the Chandra error circle of XMMU J174716.1-281048, the observed NIR band magnitudes of which suggest a low-mass star. From the available data it is inconclusive whether or not this source is the counterpart of XMMU J174716.1-281048. Future NIR spectroscopic observations will help us in determining whether or not this NIR source is related to the X-ray source. For SAX J1806.5-2215, we have very likely identified the real NIR counterpart. ",
        "introduction": "Low-mass X-ray binaries (LMXBs) are binary systems with a compact star (neutron star or black-hole) accreting matter from a low-mass star (usually K, M dwarf or red giant) while they revolve around each other. Occasionally most of these systems undergo a large increase in X-ray luminosities (by a factor more than 1000) displaying peak X-ray luminosities of 10$^{37-39}$ erg  s$^{-1}$. These X-ray  outbursts last for a few weeks to months except in some of the LMXBs in which they last up to a few years or rarely a decade as well. The sources with these long X-ray outbursts ($>$ 1 year) are also known as  quasi-persistent LMXBs. During these outbursts,  the majority of the X-rays are emitted in the inner accretion disk while the optical/infrared (OIR) light is emitted due to the reprocessed emission  in  the outer accretion disk and the surface of the companion star. In the quiescence state, the OIR light dominantly comes from the low-mass companion star and allows us to identify the companion star which is important to constrain the mass of the compact object.  During the last decade, many faint to very-faint X-ray transients have been discovered with peak X-ray luminosities of 10$^{34-36}$ erg s$^{-1}$ (\\citealt{Wijnands2006AA}, \\citealt{Degenaar2010AA}). Most of these sources are discovered\\ in the Galactic plane and a large fraction of them were discovered through the thermonuclear bursts \\citep{Cornelisse2002AA392}. These bursts have so far been observed only from the neutron star LMXBs hence probably these faint transients are also members of LMXBs. However the current instability models that explain the transient behaviour of LMXBs (e..g, \\citealt{Lasota2001}) cannot explain the low peak X-ray luminosities of these faint systems. It is possible that these sources are accreting from the stellar wind of a companion star \\citep{Pfahl2002ApJ} or they have a small accretion disk, which indicate  a small orbit that can only accommodate a very small donor star - e.g., brown dwarfs or white dwarfs - the latter are called ultra-compact X-ray binaries with extremely short orbital periods ($<$ 80 minutes; \\citealt{Nelemans2010}). To confirm these proposed scenarios, there have been efforts to identify the OIR counterpart of them (\\citealt{Kaur2011ATel3695}, \\citealt{Kaur2011ATel3268}, \\citealt{Degenaar2010MNRAS}, \\citealt{Kaur2010}, \\citealt{Kaur2009}), however most of the times, high extinction and crowded OIR fields  in the Galactic plane hamper the detection and identification of the real counterpart. In this paper we present our near-infrared (NIR) imaging observations of  two faint to very faint quasi-persistent transients to search for their counterpart.   \\subsection{XMMU J174716.1--281048}  XMMU J174716.1--281048 (hereafter XMM1747) was serendipitously discovered during  a pointed {\\it XMM-Newton} observation of the supernova remnant G0.9+0.1 performed on March 12, 2003. The source had an X-ray flux of 3.7 $\\times$ 10$^{-12}$ erg cm$^{-2}$ s$^{-1}$ in 2-10 keV energy band (\\citealt{Sidoli2003ATel147}, \\citealt{Sidoli2004}, \\citealt{DelSanto2007}). The source X-ray spectrum  was best fitted with an absorbed powerlaw model with  index $\\Gamma$ of 2.1 and a hydrogen column density N$_{\\rm H}$  of 8.9  $\\times$ 10$^{22}$ cm$^{-2}$. Being located near the Galactic center, the source distance was estimated to be  $\\sim$ 8 kpc, hence the X-ray luminosity of the source was  $\\sim$ 3 $\\times$ 10$^{34}$ erg s$^{-1}$. The source was not visible in the previous {\\it XMM-Newton} and {\\it Chandra} observations performed between September 2000 and July 2001 which implied that the quiescent X-ray luminosity of the source was $\\le$ 10$^{32}$ erg s$^{-1}$  \\citep{DelSanto2007}. The rise in flux by a factor more than 100 and  with low peak X-ray luminosity classifies the source as a very faint X-ray transient \\citep{Wijnands2006AA}.  On March 22, 2005, the {\\it INTEGRAL}/JEM-X  detected a transient IGR J17464--2811 because it exhibited a type-I X-ray burst \\citep{Brandt2006ATel970}. The position coincidence of this burst source  with the position of XMM1747, which was also active during observations made on Feb 25-26, 2005  using the {\\it XMM-Newton} satellite \\citep{Sidoli2006} indicated that both  sources are the same source \\citep{Wijnands2006ATel972}. This classifies XMM1747 as a neutron star low mass X-ray binary. On the basis of the properties of the X-ray burst observed by the {\\it INTEGRAL} satellite, it was suggested that the source was likely continuously active between 2003 and 2005 and was undergoing a long-term outburst  \\citep{DelSanto2007}. Since 2007, XMM1747 has been observed once using the  {\\it Chandra} satellite and several times using the {\\it Swift} satellite and is detected at  similar fluxes in all observations and with very similar X-ray spectral parameters  (\\citealt{Degenaar2007ATel1078}, \\citealt{Degenaar2007ATel1136}, \\citealt{Campana2009}, \\citealt{Sidoli2007ATel1174}, \\citealt{DelSanto2007ATel1207}, \\citeyear{DelSanto2007}, \\citeyear{DelSanto2009ATel2050}, \\citeyear{DelSanto2010ATel2624}, \\citeyear{DelSanto2011ATel3471}, \\citeyear{DelSanto2012ATel4099}). This suggests that the source has been in outburst for at least  8 years and hence that it belongs to the quasi-persistent X-ray transients. Another type-I X-ray burst was observed from this source on August 13, 2010 \\citep{Degenaar2011} using the {\\it Swift}/Burst Alert Telescope (BAT) .  The best known position of the source so far is RA : 17$^\\mathrm{h}$  47$^\\mathrm{m}$ 16$\\fs$16, Dec. : $-28^{\\circ}$ $10^{\\prime}$ $48${\\farcs}$0$ with an  error circle of $0\\farcs5$ \\citep{Degenaar2007ATel1136}. A possible NIR counterpart of XMM1747 was reported at this position from the observations obtained on  May 27, 2007 using instrument ANDICAM  on the 1.3-meter SMARTS telescope. The source had a {\\it H} band magnitude of 15.3 $\\pm$ 0.1 mag but was not detected in the {\\it V} and  {\\it I} band \\citep{Degenaar2007ATel1136}.   \\subsection{SAX J1806.5--2215}  SAX J1806.5--2215 (hereafter SAX1806) was discovered using the {\\it BeppoSAX}/Wide Field Cameras (WFC) through the detection of two type-I X-ray bursts observed between August 1996 to October 1997 (\\citealt{In'tzand1999}, \\citealt{Cornelisse2002}). At the time when  these X-ray bursts occured, no persistent X-ray emission was detected from the source. However, it was later found using the  {\\it RXTE}/All Sky Monitor data of the source that it was faintly active for nearly two years at this time (from early 1996 till late 1997) with the peak persistent flux of $\\simeq$ 2 $\\times$ 10$^{-10}$ erg cm$^{-2}$ s$^{-1}$ (2-10 keV; \\citealt{Cornelisse2002}). Using the unabsorbed bolometric peak flux of the brightest  burst, an upper limit of 8 kpc on source distance was obtained  \\citep{Cornelisse2002}. Using this distance, outburst accretion X-ray luminosity of the source was $\\simeq$ 2 $\\times$ 10$^{36}$ erg s$^{-1}$.  On February 22, 2011, after 12 years in quiescence SAX1806 displayed another X-ray outburst \\citep{Altamirano2011} and at the time of writing of the paper, the source is still active ( $>$ 1 year after the start of the outburst; \\citealt{Kaur2012ATel3926}). {\\it Swift}/XRT observations obtained on March 1, 2011 detected one bright source, within the  {\\it BeppoSAX} error circle of the source. The X-ray spectrum could be well fitted with a powerlaw model of $\\Gamma$ $\\sim$ 2 and N$_H$  of $\\sim$ 5.6 $\\times$ 10$^{22}$ cm$^{-2}$ \\citep{Degenaar2011ATel3202}. The unabsorbed flux of the source suggest an X-ray luminosity of  $\\sim$ 2 $\\times$ 10$^{36}$ erg s$^{-1}$ (for a  distance of 8.0 kpc;  \\citealt{Degenaar2011}) which is similar to the X-ray activity level seen during its previous outburst. A {\\it Chandra} observation  of the source provided the best known position of the source with a 0$\\farcs$6 accuracy (90\\% confidence level) at RA : 18$^\\mathrm{h}$ 06$^\\mathrm{m}$ 32$\\fs$177 and Dec : -22$^\\circ$  14$^{\\prime}$  17$\\farcs$20   \\citep{Chakrabarty2011ATel3218}. Previous {\\it Chandra} and {\\it Swift} observations of the source  obtained  between 2000-2009 did not detect the source with an upper limit on the X-ray luminosity of (0.5 - 2) $\\times$ 10$^{33}$ erg s$^{-1}$ (\\citealt{Chakrabarty2011ATel3218}, \\citealt{Degenaar2011ATel3202}, \\citealt{Campana2009}). The low peak X-ray luminosity of the source and the long outburst classify SAX1806 as a faint quasi-persistent X-ray transient.  During the X-ray outburst, a NIR counterpart was detected in our $K_{\\mathrm{s}}$ band observations obtained in March 2011 using the WHIRC instrument on the 3.5-meter WIYN telescope \\citep{Kaur2011ATel3268} and the observations will be discussed in detail in this paper. ",
        "conclusions": "In this paper, we aimed to find the possible NIR counterparts of two `faint to very faint' class  quasi-persistent X-ray binaries - XMM1747 and SAX1806, which are going through the very long X-ray outbursts. We obtained observations of XMM1747 using the PANIC instrument on Magellan telescope in 2008 while the observations of  SAX1806 were obtained in 2011 using the WHIRC instrument on 3.5-meter WIYN telescope.  We identified two NIR stars consistent with the {\\it Chandra} error circle of our source XMM1747, marked as C1 and C2 in Figure 1. We compared our observations with the UKIRT observations of the source obtained in 2006 which were also taken during the current outburst of the source and showed that the star C1 was detected with the similar magnitudes during the two epochs while the star C2 showed significant variations in all the bands. Our color-color diagram analysis  showed that star C2 is suffering a bit more extinction than most of the stars in the field. Assuming that all the NIR emission is from the companion star, the $K$ band flux of star C1 suggest a  high mass star of spectral type O or B to be the possible counterparts if it lies close to a distance of 8 kpc. Similarly, the $K$ band flux of star C2 from the Magellan observations suggest a  high mass star e.g., B type while the UKIRT observations suggest a  A-type star for given magnitudes.  The LMXBs generally show a large variations in the OIR fluxes during their X-ray outburst (reference) and none of the sources which showed type-I X-ray bursts have ever been associated with a high mass companion. Both these arguments suggest that  star C1 is not the real counterpart. However the power-law fit with an index of -2.5 suggests that it could be a disk dominated star. We note that the wavelength range for a power-law fit is small and considering that we have large uncertainty on the optical extinction, the power-law index could be misleading. Also the wind accretion from high-mass counterpart cannot be denied, and if true, XMM1747 would be the first X-ray burster from a high mass X-ray binary. From all these arguments, star C1 is likely not the real counterpart.  On the other hand, star C2 showed a significant variations during the two epochs and if the Galactic extinction as measured by \\citet{Dickey1990} represent the true extinction for this source, in which case it would be a A-type star, then it might be possible that the NIR source is associated with XMM1747.  However the fact that XMM1747 is lying in a very crowded field in the NIR bands,  which is also suggested  by detection of  two stars in the {\\it Chandra} error circle of the source show that the chance coincidence of  detection of a foreground star is very high. Also due to the high extinction towards the source, we cannot  deny that the real counterpart might not have been detected.  From the Magellan observations,  we measured the limiting magnitude of 20.0 mag in the $K$ band. Hence if the real counterpart was indeed not detected during our observations, then the observed $K$ band magnitude of the counterpart would be 20.0 mag or fainter, which indicate a  star of spectral type later than K -type, which is not unexpected for LMXBs.  For SAX1806, we detected a NIR star of magnitude 17.21 mag in $K$ band during the X-ray outburst, consistent with its {\\it Chandra} error circle. This star was not detected during the UKIRT observations which were taken during the quiescence state of the source and suggest a magnitude of 18.2 mag or fainter in $K$ band. The increase in brightness by one magnitude during the X-ray outburst as compared to the quiescence suggest the detected NIR star is the likely counterpart of SAX1806. If true, the real magnitude of the NIR counterpart would be fainter than 18.2 mag and hence suggest a companion star close to a star of spectral type F or later.  The OIR spectrum of LMXBs during outbursts show a number of double profile emission lines of H and He \\citep{Kaur2012MAXI}, indicating the presence of accretion disk. However during the quiescence state, the OIR spectrum would be dominated by the spectral signatures of the companion star. The NIR spectroscopic observations during the outburst or quiescence state can unveil if the suggested stars are the real counterparts."
    },
    "1208/1208.1292_arXiv.txt": {
        "abstract": "We study how active-region-scale flux tubes rise buoyantly from the base of the convection zone to near the solar surface by embedding a thin flux tube model in a rotating spherical shell of solar-like turbulent convection.  These toroidal flux tubes that we simulate range in magnetic field strength from 15 kG to 100 kG at initial latitudes of 1$^{\\circ} $ to 40$^{\\circ}$ in both hemispheres.  This article expands upon Weber, Fan, and Miesch (\\textit{Astrophys. J.}, 741, 11, 2011) (Article 1) with the inclusion of tubes with magnetic flux of $10^{20}$ Mx and $10^{21}$ Mx, and more simulations of the previously investigated case of $10^{22}$ Mx, sampling more convective flows than the previous article, greatly improving statistics. Observed properties of active regions are compared to properties of the simulated emerging flux tubes, including: the tilt of active regions in accordance with Joy's Law as in Article 1, and in addition the scatter of tilt angles about the Joy's Law trend, the most commonly occurring tilt angle, the rotation rate of the emerging loops with respect to the surrounding plasma, and the nature of the magnetic field at the flux tube apex.  We discuss how these diagnostic properties constrain the initial field strength of the active region flux tubes at the bottom of the solar convection zone, and suggest that flux tubes of initial magnetic field strengths of $\\ge$ 40 kG are good candidates for the progenitors of large ($10^{21}$ Mx to $10^{22}$ Mx) solar active regions, which agrees with the results from Article 1 for flux tubes of $10^{22}$ Mx.  With the addition of more magnetic flux values and more simulations, we find that for all magnetic field strengths, the emerging tubes show a positive Joy's Law trend, and that this trend does not show a statistically significant dependence on the magnetic flux. ",
        "introduction": "The toroidal magnetic field responsible for the emergence of solar active regions is believed to be generated by a dynamo mechanism at or near the base of the convection zone ({\\it{e.g.}} \\opencite{spiegel80}; \\opencite{ball82}; \\opencite{mi92}; \\opencite{gilman00}; \\opencite{char10}).  Of essential importance in the understanding of solar dynamo theory, and indeed all of solar physics, since flux emergence also regulates solar variability, is addressing how active region flux tubes rise dynamically through a turbulent convection zone. Also of interest is identifying the dynamo-generated magnetic field strength at the base of the convection zone required to produce the observed properties of solar active regions.  Valuable insights into the nature and evolution of rising magnetic loops in the solar convective envelope have been gained through the use of the thin flux tube approximation by a plethora of authors ({\\it{e.g.}} \\opencite{spruit81}; \\opencite{mi86}; \\opencite{ferriz93}; \\opencite{longcope97}; \\opencite{fan09}).  It is found that in order for these simulated flux tubes to exhibit tilt angles and latitudes of emergence that agree well with observed solar active regions, the toroidal magnetic field at the base of the convection zone needs to be in the range of  $\\approx$30 kG to $\\approx$100 kG (\\opencite{choud87}; \\opencite{sch94}; \\opencite{dsilva93}; \\opencite{cali95}).  These studies also indicate that the Coriolis force is responsible for several observed asymmetries that exist between the leading (in the direction of solar rotation) and following polarities of solar active regions, such as: the tilt angles in accordance with Joy's Law (\\opencite{dsilva93}; \\opencite{cali95}), the apparent faster proper motion of the leading polarity of an emerging active region on the solar surface (\\opencite{vandriel90}; \\opencite{mi94}; \\opencite{cali95}), and an asymmetry where the leading polarity shows a more coherent morphology (\\opencite{fan93}; \\opencite{cali95}; \\opencite{cali98}).  However, as these previous studies neglect the effects of turbulent convection on rising flux tubes, it is possible that convective downdrafts can pin portions of the flux tube down to the base of the convection zone, while helical upflows between the downdrafts can boost the rise of the flux tube (\\opencite{fan03}; \\opencite{weber2011}).  Therefore, resulting emerging loops may exhibit different properties from those produced in the absence of convection and pose an intriguing topic of study.  As of yet, the magnetic field strength at which the solar dynamo operates is not well known, nor is it directly accessible via observations.  However, solar cycle dynamo models that incorporate the Lorentz force from large scale mean fields indicate that the magnetic field strength generated and amplified at the base of the convection zone is $\\approx$15 kG, and most likely cannot exceed 30 kG (\\opencite{rempel06a}; \\opencite{rempel06b}).  Recent simulations of solar-like stars that rotate three times the current solar rate have shown that a rotating convective envelope can generate a dynamo that consists of opposite polarity magnetic wreaths in two hemispheres, which span the depth of the convection zone (\\opencite{brown2010}).  When portions of these wreaths become strong enough, $\\approx$35 kG or greater, a buoyant magnetic loop develops which then rises through the convecting fluid in which it is embedded (\\opencite{nelson2011}).  While these dynamo-producing convection simulations are not meant to reproduce the solar dynamo directly, they do demonstrate that persistent toroidal fields can coexist with convection at moderate magnetic field strengths.  In light of these studies, it is important to understand how toroidal flux tubes of weak to moderate field strengths, $\\approx$15 kG to $\\approx$50 kG, behave as they rise through a turbulent solar convective envelope.  Some studies have been performed that investigate the buoyant rise of fully three-dimensional isolated flux tubes in a turbulent convective velocity field ({\\it{e.g.}} \\opencite{fan03}; \\opencite{jouve09}).  However, due in part to the limited numerical resolution of these simulations, large values of magnetic flux must be used, which are greater than typical active region flux, in order to keep the tube from dissipating as it rises.  Recently, \\opencite{weber2011} (hereafter Article 1) incorporate a thin flux tube model in a (separately computed) rotating three-dimensional convective velocity field representative of the solar convective envelope, in an effort to study the effects of turbulent solar-like convection on the dynamic evolution of active region scale emerging flux tubes.  Although the thin flux tube model cannot capture the possible fragmentation of the flux tube and its internal magnetic structure, it does preserve the frozen-in condition of the magnetic field.  The thin flux tube model is also useful in that each simulation can be performed quickly on single processor desktops, as compared to the three-dimensional simulations, which require multi-million processor hours on massively parallel supercomputers.  As such, the thin flux tube approach is a useful platform for a parameter space study, and provides a starting point for the investigation of the effects of convective flows on flux tubes of realistic magnetic field strengths with magnetic fluxes similar to those of active regions observed on the Sun.  This article builds upon the work as presented in Article 1.  In Article 1, it was found that as the magnetic field strength increases from 15 kG to 100 kG, the dynamic evolution of the emerging flux loops changes from being convection dominated to magnetic buoyancy dominated.  For these flux tubes, the convective flow is found to reduce the rise time and the latitude of emergence through the anchoring of flux loop footpoints by downdrafts.  The addition of solar-like convection also promotes tilt angles that are consistent with the observed mean tilt of solar active regions in part because of the mean kinetic helicity of the upflows in the simulated convection. In this article, we expand the study by increasing the number of simulations performed by about 3.5 times per magnetic field strength to improve the statistics of the results, sampling different time spans of the convective flows.  Solar active regions exhibit magnetic flux in the range of $10^{20}$ Mx to $10^{22}$ Mx, which include ephemeral regions and pores to the largest scale sunspots (\\opencite{zwaan87}).  To study the dependence of flux tube evolution on magnetic flux, we also include cases with flux of $10^{20}$ Mx and $10^{21}$ Mx in this study in addition to the $10^{22}$ Mx cases. This in combination with the increased number of simulations per magnetic field strength results in a total of 5940 flux tubes to analyze.  As with Article 1, we consider flux tubes with initial latitudes of 1$^{\\circ}$ to 40$^{\\circ}$ in both the northern and southern hemispheres. Besides the observed active region properties described by Joy's Law, in this article we also consider additional observational diagnostics, including the scatter of active region tilts and sunspot rotation rates to further constrain the initial field strength of active region flux tubes at the base of the solar convection zone.  We also discuss the properties of the magnetic field at the apices of the emerging flux loops.  In comparison to Article 1, our new study confirms that with convection all flux tubes with magnetic field strengths of 15 kG -- 100 kG do exhibit a positive Joy's Law trend, and given the uncertainties that we obtain, we note no significant dependence of the flux tube tilt angle on magnetic flux.  A description of the thin flux tube model and how it is incorporated with the spherical shell of solar-like convection employed in this study is outlined in Section \\ref{sec:model}.  We highlight the results obtained from our simulations in Section \\ref{sec:results}, discussing flux tube rise times, latitudes of emergence, tilt angles, magnetic field properties, and the rotation rate of the simulated emerging loops.  A summary of the results is given in Section \\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  By embedding the thin flux tube model in a three-dimensional, turbulent, convective velocity field representing the solar convective envelope, we study how convection can influence the properties of emerging active region flux tubes.  In comparing these properties to those obtained from solar active region observations, we attempt to constrain the magnetic field strengths of dynamo-generated magnetic fields at the base of the solar convection zone. The thin flux tube approximation, although idealized, allows us to investigate active region scale flux tubes at weak to strong magnetic field strengths of 15 kG\\,--\\,100 kG under perfect frozen-in flux conditions.  We find that subjecting the thin flux tube to turbulent convective flows does indeed alter flux tube dynamics, and that it can have a significant impact on the properties of the emerging flux loop in comparison to flux tube simulations performed in the absence of a convective velocity field.  Also, the addition of convection aids the flux tube in more closely reproducing some observed properties of solar active regions.  This article extends upon the previous work of Article 1 by including flux tube simulations with additional magnetic fluxes of $10^{20}$ Mx, $10^{21}$ Mx, as well as $10^{22}$ Mx, and increasing the number of simulations performed per magnetic field strength by a factor of 3.5 in order to improve statistics such that uncertainties are reduced.  We also include more observational diagnostics, such as tilt angle scatter and sunspot rotation rate, to put further constrains on the field strength of the dynamo-generated magnetic field as the progenitors of solar active regions.  Decreasing the magnetic flux of the tube results in an increase of the drag force acting on the rising flux loop.  With convection, flux tube rise times decrease with decreasing flux for initial magnetic field strengths of 15 kG \\,--\\,60 kG because the increased drag force causes the flux tube to be more closely coupled with convection, and so they are advected by turbulent flows more strongly than the $10^{22}$ Mx case.  For all magnetic fluxes that we consider here, flux tubes are able to emerge near the equator with the aid of convective flows, which solves the previous problem of poleward slip for flux tubes of low magnetic field strengths without convective effects.  With the increased number of thin flux tube simulations, and hence improved statistics in this study, we are able to confidently say that for all magnetic field strengths of 15 kG to 100 kG, and all magnetic fluxes studied here, produce emerging loops with tilt angles that follow the Joy's Law trend.  This is an improvement upon Article 1 where the slope of the linear Joy's Law fit of the tilt angles of the emerging loops as a function of emerging latitude had too large an uncertainty for magnetic field strengths of 15 kG and 30 kG to report a definitive Joy's Law trend.  Of particular note is the fact that helical convective upflows help to drive the tilt angle of the flux tube in the appropriate Joy's law direction for both hemispheres. Including all tilt angles together for all magnetic field strengths and magnetic flux of $10^{21}$ Mx and $10^{22}$ Mx, we calculate the linear best-fit line for the tilt angle as a function of emergence latitude, and find a slope of $m_{A}=0.34$ $\\pm$ $0.02$, which overlaps with the range of $0.26$ $\\pm$ $0.05$ and $0.28$ $\\pm$ $0.06$ as suggested by white light sunspot group image from Mount Wilson and Kodikanal, respectively (\\opencite{espuig10}). Performing a fit for the tilt angle as a function of sine of the latitude for $10^{20}$ Mx, $10^{21}$ Mx, and $10^{22}$ Mx, we find a best-fit line slope of $m_{B}=22^{\\circ}$ $\\pm$ $1^{\\circ}$, which is greater than the value $15.69^{\\circ}$ $\\pm$ $0.66^{\\circ}$  obtained from white light sunspot group data by \\inlinecite{fisher95}, but less than the value of $32.1^{\\circ}$ $\\pm$ $0.7^{\\circ}$ found by \\inlinecite{stenflo12} using MDI magnetograms.  When we exclude all fields except for 40 kG -- 50 kG for all the fluxes we consider here, $m_{B}$ increases significantly to $26^{\\circ}$ $\\pm$ $2^{\\circ}$, which is closer to the value derived from magnetograms.  On the other hand, the scatter of the tilt angles around their linear Joy's Law fit line is shown to be too large for initial magnetic field strengths of 15 kG and 30 kG with fluxes of $10^{21}$ Mx and $10^{22}$ Mx, as compared to the observed value of \\inlinecite{fisher95} for white light  sunspot group images.  While the scatter of the tilt angle increases with decreasing flux (Table \\ref{tbl:scatter}), we find no statistically significant dependence of the Joy's Law trend on flux (Table \\ref{table2}), consistent with the results of \\inlinecite{stenflo12}.   We also find that the most common tilt angle produced by our study is 10.6$^{\\circ}$ for tubes with a flux of $10^{20}$ Mx, $10^{21}$ Mx, and $10^{22}$ Mx, which agrees with \\inlinecite{howard96} who finds that most tilt angles fall within the range of $7.5^{\\circ}$\\,--\\,$10^{\\circ}$ as obtained from Mount Wilson magnetogram data, although our average tilt angles are higher.  Similar to previous studies (\\opencite{cali95}; \\opencite{fan_fisher1996}; \\linebreak \\opencite{weber2011}), we find that for magnetic field strengths $\\le$ 50 kG, the leading leg of the emerging loop tends to have a larger magnetic field than the following, which may provide an explanation for the observed better cohesion of the leading polarity of an emerging active region as compared to the following polarity.  This trend of asymmetry in field strength reverses for tubes with an initial magnetic field of $\\ge$ 60 kG. However, it may be the case that the morphological asymmetry of sunspot regions is less dependent on magnetic field asymmetry, and is rather a result of the retrograde plasma flow inside the flux tube from the leading leg into the following leg as suggested by recent simulations of sunspot formation by Rempel (private communication, 2012).  If this is indeed the case, then we can not exclude magnetic field strengths of greater than 60 kG from the dynamo-generated magnetic field regime.  A study of the magnetic field of the flux tube at the top of the simulation domain suggests typical values of 500 G to 15 kG for tubes that reach $\\approx$$21$ Mm below the photosphere.  Observations show that sunspot groups tend to rotate faster than the surrounding solar surface plasma (\\opencite{howard1970}; \\linebreak \\opencite{golub1978}).  We use the apparent azimuthal motion of the center of the intersections of the emerging loop with a constant $r$ surface near the top of the simulation domain as a measure of the rotation rate of the emerging region and compare them to the average azimuthal rotation rate of the ASH convection simulation at $r=0.95R_{\\odot}$, and the surface rotation rate that we would expect assuming the surface shear layer as inferred from helioseismology.  For tubes with a flux of $10^{21}$ Mx and $10^{22}$ Mx, we find that at high emergence latitudes, the average rotation rate of the emerging loops tends to be greater than the inferred surface rotation rate for all field strengths considered. At lower latitudes, below about $35^{\\circ}$, loops with initial field strength $\\ge$60 kG tend to rotate faster than the inferred surface rate, consistent with the observed sunspot rotation rate, while loops with initial fields of about 40 kG -- 50 kG tend to rotate at a similar rate as the surface rate.  However, for initial magnetic fields below 40 kG, the rotation rate at low latitudes tends to be slower than the surface rate, contrary to observations. Thus comparison with the observed sunspot rotation rate seem to favor stronger fields as the progenitor of solar active regions.  However, because of the limitations of our model, we recognize that there remain large uncertainties in our results of the azimuthal motion of emerging loops and how they relate to the observed sunspot rotations.  These limitations arise because our simulations stop before the tube enters the solar surface shear layer, and because we do not address the deformation and fragmentation of the flux tubes which may result in a stronger and more complex coupling of the magnetic fields with the fluid motion, especially in the upper layers of the convection zone.  Overall, the results in this study suggest that the initial field strength of active region progenitor flux tubes needs to be sufficiently large, probably $\\ge$ 40 kG, in order for them to satisfy the Joy's Law trend for mean tilt angles as well as the observed amount of scatter of the tilt angles about the mean Joy's Law behavior.  Weaker magnetic fields tend to produce too large a scatter to be consistent with the observed results.  Mid-field strengths of 40 kG\\,--\\,50 kG, which take the longest to rise, do the best job at matching magnetogram observations of Joy's Law dependence.  Although only 50 kG or greater field strengths can rotate at or faster than the solar surface rate.  So, according to our thin flux tube approach, magnetic field values need to be of moderate to large field strengths for tubes with fluxes of $10^{21}$ Mx and $10^{22}$ Mx to produce sunspot rotation behavior.  The addition of multiple magnetic flux values to this study allows us to perform a more comprehensive study on the dependence of flux tube evolution with regard to magnetic flux.  We find, that with convective effects, the tube rise time substantially decreases with decreasing flux.  There also is no statistically significant dependence on the Joy's Law trend in relation to the value of the magnetic flux, although the spread of the tilt angles about the Joy's Law trend does increase as flux decreases.  Convective effects are important to flux tube development at these field strengths and magnetic flux values, and should be incorporated into future studies.  With this new study, all magnetic field strengths now show a positive Joy's Law trend, and we are now learning that the magnetic field asymmetry of sunspots could be due to a retrograde flow of plasma along the emerging flux tube.  As a result, the estimate of the required $\\ge$ 40 kG magnetic field strength is based here mainly on the tilt angle scatter and rotation rates.  However, identifying the rotation rate of the emerging flux tube, in particular, is an aspect where the thin flux tube approach may be lacking in its ability to adequately capture all of the physical processes involved.  We suggest that this aspect of flux tube evolution would be a useful topic for further high-resolution three-dimensional MHD simulations of global flux emergence. In addition, more of these simulations are needed using realistic active region flux tube parameters to further constrain the magnetic field strength generated by the solar dynamo based on observational diagnostics.        \\begin{acks} This work is supported by NASA SHP grant NNX10AB81G to the National Center for Atmospheric Research (NCAR).  NCAR is sponsored by the National Science Foundation.  We would like to thank Nick Featherstone for reading our manuscript as our internal reviewer, and for offering helpful comments.  Also, we would like to thank our referee for offering a critical review of our manuscript, which contributed to the production of a more substantial article. \\end{acks}"
    },
    "1208/1208.4362_arXiv.txt": {
        "abstract": "We measure the evolution of the specific star formation rate (${\\rm sSFR = SFR} / M_{\\rm stellar}$) between redshift 4 and 6 to investigate the previous reports of ``constant'' sSFR at $z>2$, which has been in significant tension with theoretical expectations. We obtain photometry on a large sample of galaxies at $z\\sim4\\mbox{--}6$ located in the GOODS South field that have high quality optical and IR imaging from \\emph{HST} and \\emph{Spitzer}. We have derived stellar masses and star formation rates (SFRs) through stellar population modeling of the rest-frame UV and optical spectral energy distributions (SEDs).  We estimate the dust extinction from the observed UV colors. In the SED fitting process we have studied the effects of assuming a star formation history (SFH) both with constant SFR and one where the SFR rises exponentially with time. The latter SFH is chosen to match the observed evolution of the UV luminosity function. We find that neither the mean SFRs nor the mean stellar masses change significantly when the rising SFR (RSF) model is assumed instead of the constant SFR model. We estimate the sSFR at $z>4$ for galaxies in two stellar mass bins centered at 1 and $5\\times10^9\\msun$. We find that the galaxies in the lower mass bin have very similar sSFRs to the more massive ones (within $\\sim 0.1\\,$dex). When focusing on galaxies with $M_{\\rm stellar}\\sim5\\times 10^{9}\\msun$, we find that the sSFR evolves weakly with redshift (${\\rm sSFR(z) \\propto(1+z)^{0.6\\pm0.1}}$), consistent with previous results and with recent estimates of the sSFR at $z\\sim2\\mbox{--}3$ using similar assumptions. We have also investigated the impact of optical emission lines on our results. We estimate that the contribution of emission lines to the rest-frame optical fluxes is only modest at $z\\sim4$ and 5 but it could reach $\\sim50$\\% at $z\\sim6$. When emission lines of this strength are taken into account, the sSFR shows somewhat higher values at high redshifts, according to the relation ${\\rm sSFR(z)\\propto (1+z)^{1.0\\pm0.1}}$, i.e., the best fit evolution shows a value $\\sim2.3\\times$ higher at $z\\sim6$ than at $z\\sim2$.  However, the observed evolution is substantially weaker than that found at $z<2$ or that expected from current models (which corresponds to ${\\rm sSFR(z)\\propto (1+z)^{2.5}}$). ",
        "introduction": "Large samples of Lyman break galaxies (LBGs) have allowed for the study of the properties of high redshift galaxies up to $z\\sim8$ \\citep[e.g., ][]{bouw12,bouw11b, dunl12a, fink11, fink10, gonz10, gonz11, labb10, labb10a, lee12, mclu10, mclu11, oesc10, oesc12, papo11, scha10, star09, wilk11}.  These samples are the result of large investments on high quality \\emph{Hubble Space Telescope} (\\emph{HST}) data over deep fields like GOODS \\citep{giav04} that have rich complementary multi-wavelength coverage.  Several studies have explored the observed UV and optical colors of these galaxies. Through the standard technique of synthetic stellar population modeling of the observed spectral energy distributions (SEDs), the physical properties of these galaxies, such as the stellar mass (\\mstar), star formation rate (SFR), dust attenuation, age, etc, have also been explored \\citep[e.g., ][]{yan06, eyle05, star09, gonz10, gonz11, papo11, lee12, bouw09, bouw12}.  These studies have shown that, at rest-frame UV and optical wavelengths where LBGs are  more amenable to observations, the SEDs show very similar colors with weak trends of bluer colors as a function of decreasing UV luminosity and increasing redshift (\\citealt{bouw09, bouw12, wilk11, fink11}; but see also \\citealt{dunl12a}). As a direct consequence of this, the physical properties estimated through SED fitting, using simple models, are also remarkably similar.  Particularly intriguing are the results that indicate that the specific star formation rate (sSFR) of sources with a given \\mstar~remains approximately constant from $z\\sim2$ to $z\\sim7$ \\citep{star09, gonz10, mclu11}. In particular, \\citet{gonz10} shows that for sources with $M_{\\rm stellar} = 5\\times10^9\\,\\msun$, the sSFR shows no evidence for significant evolution (${\\rm sSFR }\\sim 2\\, \\gyr^{-1}$) from $z\\sim7$ to $z\\sim2$. Such a result is at odds with the fairly generic theoretical expectation that the sSFR should decrease monotonically with cosmic time \\citep[e.g., ][]{wein11, khoc11}.  Although suggestive, much of the early work pointing toward a relatively constant sSFR was limited in many important aspects and based on a number of simplifying assumptions. For example, early work assumed zero dust extinction for the SFRs of $z>4$ galaxies. This was motivated by the very blue UV colors observed in the galaxy SEDs, which have subsequently been measured more accurately thanks to the better wavelength coverage provided by \\emph{HST}/WFC3 observations \\citep[e.g., ][]{bouw12}.  Early works also adopted exponentially declining or constant star formation histories (SFHs). This is in apparent contradiction with the evolution of the UV luminosity function (LF), which suggests that rising SFHs are a better match for the evolution of galaxies in early cosmic times \\citep[e.g., ][]{papo11}. Finally, much of the early work did not correct for the effect of optical emission lines on the estimate of stellar masses (\\mstar), and hence on the sSFRs.  Several recent studies have attempted to redress some of these shortcomings. \\citet{bouw12}, for example, used new measures of the dust extinction (based on UV colors) to correct the earlier sSFR estimates at $z\\sim4\\mbox{--}7$. \\citet{scha10} considered the possible effects of the optical emission lines on the stellar masses and sSFRs derived for high redshift galaxies, and this effort has been considerably extended in \\citet[see also \\citealt{curt12}]{de-b12}.  In this paper we attempt to bring more consistency in the exploration of the aforementioned issues, exploring the effect of dust reddening in the UV colors, SFRs, and \\mstar~measurements, as well as the effects of choosing a SFH that better matches the evolution of the UV LF, simultaneously. The goal is to make estimates for the physical properties of $z>4$ LBG that use empirically motivated assumptions that better match a larger range of observations. We plan to investigate whether this new set of assumptions still shows an approximate plateau in the sSFR at $z>2$ or shows evidence for an increase toward high redshift as predicted by theory.  We now provide a brief plan for this paper. In Section 2, we briefly describe the observational data and selection criteria used. In Section 3, we describe our approach to stellar population modeling, detailing the specific assumptions that we make and the effects these assumptions have on the physical properties we derive for LBGs. In Section 4 we present the new measurements of the sSFR at high-redshift. We discuss the results in 5, and summarize our findings in Section 6.  Throughout, we use a ($H_0,\\,\\Omega_M,\\,\\Omega_\\Lambda$) = ($70\\,\\rm{km ~s^{-1}}$, 0.3, 0.7) cosmology when necessary and we quote all magnitudes in the AB system \\citep{oke83}. ",
        "conclusions": "Sample Summary.} \\tablehead{ & \\colhead{$z\\sim4$} & \\colhead{$z\\sim5$} & \\colhead{$z\\sim6$} } \\startdata ERS    &  270 (524)  &  77 (123)   &  29 (36)\\\\ HUDF   &  137 (205)  &  41 (55)    &  54 (59)\\\\ CANDELS &          &             &  143 (182)\\\\ \\tableline\\\\ TOTAL  &  407 (729) &  118 (178)  &  226 (277)\\\\ \\enddata \\tablecomments{ Number of sources in our $z\\sim4$, $z\\sim5$, and $z\\sim6$ samples. IRAC photometry of these sources requires fitting and subtraction of the flux from surrounding foreground neighbors.  This is not possible in all the cases. The table shows the number of sources in our samples with clean IRAC photometry and the total number of sources in these samples in parenthesis. To improve the number statistics at $z\\sim6$ we have extended the sample by including the GOODS-S CANDELS data. } \\end{deluxetable} }  \\newcommand{\\tableSSFR}[1]% { \\begin{deluxetable*}{lccccccccc}[#1] \\tablecaption{\\label{tbl:SSFRresults} Current Estimates of the mean Specific Star Formation Rate for $z\\sim4\\mbox{--}6$ Galaxies} \\tablehead{ & \\multicolumn{3}{c}{$z\\sim4$} & \\multicolumn{3}{c}{$z\\sim5$} & \\multicolumn{3}{c}{$z\\sim6$} \\\\[0.4em] \\colhead{Model} & \\colhead{$N_{\\rm bin}$} & \\colhead{$N_{\\rm rej}$} & \\colhead{$\\log_{10}({\\rm sSFR/Gyr^{-1}})$}& \\colhead{$N_{\\rm bin}$} & \\colhead{$N_{\\rm rej}$} & \\colhead{$\\log_{10}({\\rm sSFR/Gyr^{-1}})$}& \\colhead{$N_{\\rm bin}$} & \\colhead{$N_{\\rm rej}$} & \\colhead{$\\log_{10}({\\rm sSFR/Gyr^{-1}})$} } \\startdata \\multicolumn{7}{c}{$M_{\\rm stellar} = 5\\times10^9\\ M_\\odot$}\\\\[0.3em] CSF no emission lines    & 54 & 2 & 0.50($\\pm0.05$)     & 18 & 0 & 0.57($\\pm0.08$)     & 73 & 3 & 0.42($\\pm0.04$)     \\\\ CSF with emission lines  & 51 & 2 & 0.55($\\pm0.05$)     & 15 & 0 & 0.60($\\pm0.06$)     & 48 & 5 & 0.54($\\pm0.05$)     \\\\ RSF no emission lines    & 67 & 0 & 0.49($\\pm0.03$)     & 21 & 0 & 0.48($\\pm0.04$)     & 78 & 4 & 0.60($\\pm0.03$)     \\\\ RSF with emission lines  & 68 & 0 & 0.51($\\pm0.03$)     & 19 & 0 & 0.50($\\pm0.04$)     & 71 & 5 & 0.65($\\pm0.04$)     \\\\[.3em] \\tableline\\\\[-0.7em] \\multicolumn{7}{c}{$M_{\\rm stellar} = 1\\times10^9\\ M_\\odot$}\\\\[0.3em] CSF no emission lines    & 99 & 6 & 0.60($\\pm0.04$)     & 32 & 4 & 0.66($\\pm0.07$)     & 57 & 13 &0.79($\\pm0.08$)     \\\\ CSF with emission lines  & 100& 6 & 0.61($\\pm0.04$)     & 35 & 4 & 0.68($\\pm0.07$)     & 85 & 18& 0.80($\\pm0.06$)     \\\\ RSF no emission lines    & 117& 2 & 0.58($\\pm0.03$)     & 36 & 2 & 0.67($\\pm0.06$)     & 74 & 2 & 0.61($\\pm0.04$)     \\\\ RSF with emission lines  & 117& 2 & 0.61($\\pm0.03$)     & 40 & 2 & 0.73($\\pm0.07$)     & 81 & 3 & 0.73($\\pm0.05$) \\enddata \\tablecomments{ Mean values of the estimated ${\\rm \\log_{10}(sSFR/Gyr^{-1})}$ for our samples using different model assumptions. The sSFR was estimated for two stellar mass bins centered at $\\log_{10}(M_{\\rm stellar}/M_\\odot)=9.7$ and 9.0. The width of each bin is $\\pm0.3$ dex. The $N_{\\rm bin}$ column indicates how many sources fall in each bin at each redshift for a given model.  Extreme values of the sSFR (${\\rm \\sim100~Gyr^{-1}}$, which correspond to minimum age models) were rejected before taking the mean. This latter choice does not make a significant difference (see Figure \\ref{fig:SSFRcomparison}). } \\end{deluxetable*} }    \\begin{document} \\title{Slow Evolution of the Specific Star Formation Rate at $z>2$: The Impact of Dust, Emission Lines, and A Rising Star Formation History }  \\author{ Valentino Gonz\\'alez\\altaffilmark{1,2}, Rychard Bouwens\\altaffilmark{3}, Garth Illingworth\\altaffilmark{1}, Ivo Labb\\'e\\altaffilmark{3}, Pascal Oesch\\altaffilmark{1,4}, Marijn Franx\\altaffilmark{3}, \\textsc{and} Dan Magee\\altaffilmark{1} } \\altaffiltext{1}{Astronomy Department, University of California, Santa Cruz, CA 95064} \\altaffiltext{2}{Department of Physics and Astronomy, University of California, Riverside, CA 92521} \\altaffiltext{3}{Leiden Observatory, Leiden University , NL-2300 RA Leiden, The Netherlands} \\altaffiltext{4}{Hubble Fellow}  \\begin{abstract}  We measure the evolution of the specific star formation rate (${\\rm sSFR = SFR} / M_{\\rm stellar}$) between redshift 4 and 6 to investigate the previous reports of ``constant'' sSFR at $z>2$, which has been in significant tension with theoretical expectations. We obtain photometry on a large sample of galaxies at $z\\sim4\\mbox{--}6$ located in the GOODS South field that have high quality optical and IR imaging from \\emph{HST} and \\emph{Spitzer}. We have derived stellar masses and star formation rates (SFRs) through stellar population modeling of the rest-frame UV and optical spectral energy distributions (SEDs).  We estimate the dust extinction from the observed UV colors. In the SED fitting process we have studied the effects of assuming a star formation history (SFH) both with constant SFR and one where the SFR rises exponentially with time. The latter SFH is chosen to match the observed evolution of the UV luminosity function. We find that neither the mean SFRs nor the mean stellar masses change significantly when the rising SFR (RSF) model is assumed instead of the constant SFR model. We estimate the sSFR at $z>4$ for galaxies in two stellar mass bins centered at 1 and $5\\times10^9\\msun$. We find that the galaxies in the lower mass bin have very similar sSFRs to the more massive ones (within $\\sim 0.1\\,$dex). When focusing on galaxies with $M_{\\rm stellar}\\sim5\\times 10^{9}\\msun$, we find that the sSFR evolves weakly with redshift (${\\rm sSFR(z) \\propto(1+z)^{0.6\\pm0.1}}$), consistent with previous results and with recent estimates of the sSFR at $z\\sim2\\mbox{--}3$ using similar assumptions. We have also investigated the impact of optical emission lines on our results. We estimate that the contribution of emission lines to the rest-frame optical fluxes is only modest at $z\\sim4$ and 5 but it could reach $\\sim50$\\% at $z\\sim6$. When emission lines of this strength are taken into account, the sSFR shows somewhat higher values at high redshifts, according to the relation ${\\rm sSFR(z)\\propto (1+z)^{1.0\\pm0.1}}$, i.e., the best fit evolution shows a value $\\sim2.3\\times$ higher at $z\\sim6$ than at $z\\sim2$.  However, the observed evolution is substantially weaker than that found at $z<2$ or that expected from current models (which corresponds to ${\\rm sSFR(z)\\propto (1+z)^{2.5}}$).  \\end{abstract}  \\keywords{galaxies: evolution --- galaxies: high-redshift}    Many of the earliest studies of the sSFR at high-redshift had suggested that the sSFR did not evolve strongly at $z\\gtrsim2$ \\citep{gonz10, star09, mclu11}. This is very different from the behavior observed at lower redshifts and also from the expectations of theoretical models that consistently predict a sSFR that declines monotonically with cosmic time when halos of a constant mass are studied \\citep[e.g., ][]{wein11, bouc10, dave08}. Most earlier studies, however, did not determine the SFRs and \\mstar~in a way that was clearly self-consistent, did not adopt SFHs which are consistent with the evolution of the UV LF, and also did not account for the effects of nebular emission lines in the stellar population modeling.  More recently, however, there have been some efforts to redress these shortcomings. For example, \\citet{bouw12} focus on correcting previous sSFR determinations to reflect the latest estimates of dust extinction. Another example is the work of \\citet{de-b12} where there is an effort to correct for the impact of the emission lines on the inferred stellar masses. \\citet{star13} also explores the effect of emission lines using a modeling technique very similar to the one presented here. In general, all these recent studies, including the present work, have suggested somewhat higher sSFRs at $z\\gtrsim4$. It is worth noting, though, that despite the growing consensus that the SFRs at $z\\gtrsim4$ increase, there is fairly large difference in the magnitude of evolution derived. In the following we compare to a couple of recent results with the goal of highlighting the main analysis differences that yield to the different results.  \\subsection{Comparison to \\citet{de-b12}}  In \\citet{de-b12} the authors find an order of magnitude larger sSFRs at $z>4$ than is found at $z\\sim2$, the sSFRs we find at $z\\gtrsim4$ are only larger than the $z\\sim2$ values by factors of $\\sim1.5\\mbox{--}2.0$. What are the reasons for these differences? Largely, this is the result of the different assumptions used in modeling the observed photometry. Differences in assumptions can have an impact on the derived sSFRs due to degeneracies between different model parameters.  While we adopt the simplifying assumption that the observed UV continuum slope can be used to estimate the dust extinction in high-redshift galaxies, \\citet{de-b12} do not impose any constraints on the dust, metallicity, mass, or age in modeling the photometry of high-redshift galaxies. By not restricting the model parameter space through various simplifying assumptions, \\citet{de-b12} observe considerable scatter in the properties of many of their sources and also find a large fraction of sources to have very young ages (a result which we consider probably unphysical and may result in \\citealt{de-b12} reporting very high sSFRs). On the other hand, with our approach, we use the UV continuum slope to set the dust extinction, effectively breaking many of the model degeneracies.  While one can debate which approach yields the most accurate results, we prefer our approach due to the strong evidence at both $z\\sim0$ and $z\\sim2$ that the UV continuum slope correlates on average with the observed dust extinction \\citep{meur99, burg05, over11, dadd07, redd12}.  \\sSFR{figure*}{ht!}{0.95}   \\subsection{Comparison to \\citet{star13}}  \\muvmstar{figure*}{t}{0.85}  In a recent study \\citet{star13} explores the impact of rest-frame optical emission lines on the stellar masses and sSFR derived through SED fitting at $z\\gtrsim4$. Their analysis is similar in many regards to the one presented here but they find a sSFR evolution that is much faster with redshift, in agreement with theoretical expectations.  As we discuss below, it appears that the reason for this difference is that the consideration of the effect of the $M/L$ scatter turns out to play an important role.  The SED modeling assumptions used in \\citet{star13} are very similar to the ones we have used here, and the resulting SFR and stellar mass estimates for individual sources are in good agreement, as can be seen in Figure \\ref{fig:muvmstar}. Only sources that are detected in IRAC ($>2\\,\\sigma$) are shown in the figure (background points). The large open squares correspond to our estimates for the median SEDs stacked by UV luminosity from \\citet{gonz11} and the dashed lines correspond to the trends derived by \\citet{star13} based on the same stacked SEDs.  This figure makes it clear that there is broad overall agreement in the basic properties derived.  The way in which the average sSFR is determined at a given redshift, however, causes important differences in the conclusions. The results presented in the \\citet{star13} work are based on the best fit $\\log(M_{\\rm stellar})\\mbox{--}M_{\\rm UV}$ relation derived from the stacked SEDs. These stacked SEDs are binned according to their UV luminosity. Their average sSFRs at a given redshift, then, correspond to UV luminosity binned averages. This is similar to the result by \\citet{bouw12}, who only apply an improved dust correction to the luminosity binned results but \\citet{star13} also include mass corrections due to the effect of emission lines.  In the present analysis we estimate the sSFR in bins of stellar mass. Within a given bin there will be galaxies with lower than average luminosities but high $M/L$ ratios, as well as bright galaxies with low $M/L$ ratios.  Since fainter galaxies are more numerous, the distribution of $M/L_{\\rm UV}$ ratios within the bin is skewed to high $M/L$ ratios (equivalent to lower sSFRs, see large filled squares in Figure \\ref{fig:muvmstar}).  \\citet{star13} recognize and discuss the differences that the $M/L$ ratio distribution at a given mass cause on the mean sSFR \\citep[see also][]{redd12}. They report that, at $z\\sim4$, the mass binned average sSFR could be $2.8\\times$ lower if a symmetric scatter of 0.5\\,dex (the observed scatter reported in \\citealt{gonz12}) was assumed instead of no scatter. Nevertheless, they report their results assuming zero scatter in $M/L$ and suggest that the above effect could be offset by possibly higher UV luminosity to SFR conversion factors at higher redshift. Meanwhile, in our estimates we consider the \\emph{observed} $M/L$ distribution at a given mass. Neither approach is perfectly correct, but this explains why our estimates of the sSFRs are lower.  Ideally, we should use the intrinsic distribution of $M/L$ ratios, as opposed to the observed distribution which is broader due to modeling and observational uncertainties. A reliable estimate of the intrinsic scatter in M/L ratios is out of the scope of this paper given the absence of rest-frame optical spectroscopy, and will likely have to wait until JWST. Here we just note that our results reported in Table \\ref{tbl:SSFRresults} are consistent with those of \\citet{star13} if a symmetric intrinsic scatter of 0.3\\,dex in $\\log(M_{\\rm stellar})$ is added to their mean relation at $z\\sim4$ (0.2 dex at $z\\sim6$).  It is interesting to note that in the works of \\citet{de-b12}, \\citet{star13}, as well as our study, the sSFR evolution from $z\\sim4$ to $z\\sim5$ is quite small (for a given set of assumptions, in particular, for SFH and emission line strengths). As shown in Figure \\ref{fig:mstarmstar}, at these redshifts emission lines have only a weak effect on the stellar masses and SFRs.  Differences in the sSFR appear only at $z>5$, which is also when emission lines can have a bigger impact on the stellar population modeling. Furthermore, there is only $\\sim240$ Myr between $z\\sim5$ and $z\\sim6$. Even though it is possible that the differences in sSFR arising between $z\\sim6$ and $z\\sim5$ are real, it seems surprising that significant changes manifest themselves at the same redshift where emission lines potentially play a large role in the SED modeling (especially considering how uncertain the assumed EWs are).  We conclude that the sSFR \\emph{at a constant stellar mass} changes only weakly with cosmic time when full consideration is given to the scatter in $M/L$ -- with or without the consideration of emission lines (Figures \\ref{fig:SSFRcomparison} and \\ref{fig:sSFR}). A fit to the \\emph{maximal} emission line estimates of the sSFR, which includes the values at $z\\sim2\\mbox{--}3$ reported in the literature \\citep{dadd07, redd12}, yields a best-fit to the evolution of the ${\\rm sSFR} \\propto (1+z)^{1.0\\pm0.1}$, indicating that the best fit sSFR is higher by $\\sim2.3\\times$ at $z\\sim6$ than at $z\\sim2$ (compared to the factor $\\sim8.3\\times$ predicted in most simulations). When emission lines are not included in the models (as in our CSF model), the exponent is lower: 0.6$\\pm0.1$. The evolution we derive at $z\\gtrsim2$ is less than expected from the lower-$z$ trends and is not consistent with the predictions from numerical simulations \\citep[e.g., ][]{neis08, wein11}."
    },
    "1208/1208.4681_arXiv.txt": {
        "abstract": "The growth rate of matter density perturbations has been measured from redshift-space distortion (RSD) in the galaxy power spectrum.  We constrain the model parameter space for representative modified gravity models to explain the dark energy problem by using the recent data of $f_m(z) \\sigma_8(z)$ at the redshifts $z = $ 0.06--0.8 measured by WiggleZ, SDSS LRG, BOSS, and 6dFGRS. We first test the Hu-Sawicki's $f(R)$ dark energy model, and find that only the parameter region close to the standard $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) model is allowed ($\\lambda > $ 12 and 5 for $n = $ 1.5 and 2, respectively, at 95\\% CL).  We then investigate the covariant Galileon model with a de Sitter attractor and show that the parameter space consistent with the background expansion history is excluded by the RSD data at more than $8 \\sigma$ because of the too large growth rate predicted by the theory. Finally, we consider the extended Galileon scenario, and we find that, in contrast to the covariant Galileon, there is a model parameter space for a tracker solution that is consistent with the RSD data within a $2\\sigma$ level. ",
        "introduction": "The observational support for the existence of dark energy \\cite{SN98,CMB03,LSS04} has motivated the idea that the gravitational law may be modified from General Relativity (GR) at large distances to realize the late-time cosmic acceleration (see Refs.~\\cite{Sotiriou,DT10,fRde,Clifton} for reviews). In this vein many modified gravity models of dark energy have been proposed---including those based on $f(R)$ gravity \\cite{fR}, scalar-tensor theories \\cite{stensor}, the Dvali-Gabadadze-Porrati (DGP) braneworld scenario \\cite{DGP}, Galileons \\cite{Nicolis,Deffayet}, and so on.  Measuring the growth rate of large scale structures in the Universe gives a strong test for the modified gravity scenario, because modified gravity models generally predict different growth rates from that in the standard $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) model. In fact the equations of linear matter density perturbations have been derived for a number of modified gravity models---including $f(R)$ gravity \\cite{fRper,Tsujikawa07}, the DGP model \\cite{DGPper}, and Galileons \\cite{Kase}.  Recently De Felice {\\it et al.} \\cite{DKT} derived the full perturbation equations as well as the effective gravitational coupling to non-relativistic matter in the most general scalar-tensor theories in 4 dimensions \\cite{Horndeski,DGSZ,KYY11} (which cover most of the single-field dark energy models proposed in the literature).  One of the methods to measure the cosmic growth rate is redshift-space distortion (RSD) that appears in clustering pattern of galaxies in galaxy redshift surveys because of radial peculiar velocities. RSD on large and linear scales reflects the velocity of inward collapse motion of large scale structure, which is directly related to the evolutionary speed of matter overdensity perturbations simply from the mass conservation \\cite{Kaiser}.  Recent galaxy redshift surveys have provided measurements of the growth rate $f_m(z)$ or $f_m(z) \\sigma_8(z)$ as a function of redshift up to $z \\sim$ 1 \\cite{Tegmark04,Percival04,Tegmark06,Yamamoto,Guzzo08,Blake,Samushia11,Reid12,Beutler12,Samushia12,Hudson,Jailson}, where $f_m = d\\ln \\delta_m/d\\ln a$, $\\delta_m$ is the fractional over-density of non-relativistic matter, $a$ is the scale factor of the Universe, and $\\sigma_8$ is the rms amplitude of over-density at the comoving 8\\,$h^{-1}$ Mpc scale ($h$ is the normalized Hubble parameter $H_0=100\\,h$~km\\,sec$^{-1}$Mpc$^{-1}$). In this paper, we constrain some of the modified gravity models proposed to solve the dark energy problem, by using the latest RSD data of 2dFGRS \\cite{Percival04}, WiggleZ \\cite{Blake}, SDSS LRG \\cite{Samushia11}, BOSS \\cite{Reid12}, and 6dFGRS \\cite{Beutler12}.  In this paper, we choose the $f(R)$ gravity and Galileon models as the representative theories of modified gravity.  A general difficulty of dark energy models based on modified gravity is the emergence of ``the fifth force'' that can violate the constraints from local gravity experiments.  There are a number of mechanisms to suppress the propagation of the fifth force in local regions: (i) chameleon mechanism \\cite{chame}, (ii) Vainshtein mechanism \\cite{Vainshtein}, and (iii) the symmetron mechanism \\cite{symmetrons,BBDS}. The symmetron mechanism is irrelevant to dark energy because the energy scale of its simplest potential is too small to be used for the late-time cosmic acceleration\\footnote{Unless the field potential is carefully designed, this problem even persists for the chameleon field \\cite{Brachame}.}. The $f(R)$ gravity and Galileons are representative models to explain the acceleration, avoiding the fifth force problem by mechanisms (i) and (ii), respectively (see Sec.~\\ref{modelsec}, where we give a brief review of these theories).  In $f(R)$ gravity several authors put observational bounds on the parameter $B=(F_{,R}\\dot{R}/F)(H/\\dot{H})$ by assuming the $\\Lambda$CDM background \\cite{Song07,Peiris,Lombriser}, where $F=\\partial f/\\partial R$, $F_{,R}=\\partial F/\\partial R$, and $H$ is the Hubble parameter (see also Refs.~\\cite{fRobsercon} for related works).  At the level of perturbations the parameter $B$ characterizes the deviation from the $\\Lambda$CDM model.  The joint data analysis of cluster abundance, the cosmic microwave background (CMB), and other observations shows that the value of $B$ today is constrained to be $B_0<1.1 \\times 10^{-3}$ at the 95~\\% confidence level (CL) \\cite{Lombriser}. The matter and the velocity power spectra were also computed with $N$-body simulations for the Hu-Sawicki model (a popular and viable $f(R)$ dark energy model, see Eq.~(\\ref{fRmo}) below) with $n=1/2$ \\cite{Jennings}. However, those past works did not place explicit constraints on the viable parameter space of the Hu-Sawicki model from the RSD data.  Recently the covariant Galileon dark energy model (whose cosmological dynamics was studied in Refs.~\\cite{GS10,DTPRL,DTPRD}) was confronted with observations \\cite{ApplebyLinder} by using the RSD data of WiggleZ and BOSS as well as the data of Supernovae Ia (SN Ia), CMB, and Baryon Acoustic Oscillations (BAO).  It was found that this model is severely disfavored over the $\\Lambda$CDM.  In this paper we show that the covariant Galileon is indeed excluded at more than $8 \\sigma$ CL by using the most recent RSD data of SDSS LRG and 6dFGRS in addition to the WiggleZ and BOSS data. We then further consider the extended Galileon scenario in which a tracker solution with an arbitrary constant dark energy equation of state smaller than $-1$ is realized during the matter era. Unlike the covariant Galileon we show that there are some viable parameter spaces compatible with the current observational data of RSD as well as SN Ia, CMB, and BAO.  This paper is organized as follows.  In Sec.~\\ref{modelsec} we review the basic properties of dark energy models based on $f(R)$ gravity and Galileons.  In Sec.~\\ref{persec} we study the evolution of the growth rate of matter perturbations in these models. In Sec.~\\ref{rsdconsec} we put observational constraints on the parameters of each model, and Sec.~\\ref{consec} is devoted to conclusions. ",
        "conclusions": "\\label{consec}  In this paper we have placed observational constraints on dark energy models based on $f(R)$ gravity, covariant Galileon, and extended Galileon from the latest data of galaxy redshift surveys (WiggleZ, SDSS LRG, BOSS, and 6dFGRS).  In these models the General Relativistic behavior can be recovered at short distances under the chameleon mechanism or the Vainshtein mechanism.  On scales relevant to large-scale structures the modification of gravity manifests itself for the growth rate of matter density perturbations. This growth rate is related to the peculiar velocities of galaxies, which can be constrained from the redshift-space distortions.  As an explicit example we studied the $f(R)$ model (\\ref{fRmo}) proposed by Hu and Sawicki, which is general enough to cover basic properties of $f(R)$ dark energy models.  At the level of perturbations this model mimics the $\\Lambda$CDM for the redshift $z$ larger than the critical value $z_c$ given by Eq.~(\\ref{zc}), but the deviation appears for $z<z_c$.  For smaller values of $n$ and $\\lambda$ the transition redshift $z_c$ gets larger, so that the growth of matter perturbations is more significant at late times.  We placed observational bounds on the Hu-Sawicki $f(R)$ model in the $(n, \\lambda)$ plane.  Since the growth rate in this model depends on the scale of density perturbations, we derive constraints assuming three different wavenumbers of $k^{-1}=10, \\ 30,$ and $60\\, h^{-1}$~Mpc which are relevant to the galaxy redshift surveys considered here. As we see in Fig.~\\ref{fig6} (which corresponds to $\\sigma_8(z=0) = 0.811$ and $\\Omega_{\\rm DE}^{(0)} = 0.73$), the constraints on model parameters tend to be tighter for smaller scales.  For $k^{-1}=60 \\, h^{-1}$~Mpc the models with $n<2$ and $\\lambda<3$ are outside the 2$\\sigma$ observational contour, whereas for $k^{-1}=10 \\, h^{-1}$~Mpc the models with $n<2$ and $\\lambda<8$ are excluded at the 2$\\sigma$ CL.  These results show that the recent RSD data do not favor the large deviation of $f_{m}\\sigma_8$ from that in the $\\Lambda$CDM, whose property can be confirmed in Figs.~\\ref{fig1} and \\ref{fig2}.  For covariant Galileon there is a tracker solution along which the equation of state $w_{\\rm DE}$ of dark energy changes from $-2$ (matter era) to $-1$ (de Sitter era). The likelihood analysis of SN Ia, CMB, and BAO shows that only the late-time tracking solution is allowed from the data. Using the parameter space constrained from the background cosmology, we found that the covariant Galileon model is excluded at more than $8 \\sigma$ CL for $\\sigma_8 (z=0)$ ranging in the region $0.75<\\sigma_8 (z=0)<0.85$. This is associated with the fact that the cosmic growth rate for the covariant Galileon is much larger than that in the $\\Lambda$CDM, as confirmed in Fig.~\\ref{fig3}. In the range $10\\,h^{-1}$~Mpc $\\lesssim\\,k^{-1} \\lesssim$ $60\\,h^{-1}$~Mpc the evolution of $f_m \\sigma_8$ is practically independent of the scales.  In the extended Galileon scenario the equation of state for the tracker is given by $w_{\\rm DE}=-1-s$ during the matter era, where $s$ is a positive constant. For the model with $s=0.2$ we found that there is a parameter space which is inside the $2\\sigma$ observational contour. As we see in Fig.~\\ref{fig8}, the models with larger values of $\\alpha$ and $\\beta$ tend to be favored from the RSD data.  According to the latest data release of Planck, the dark energy density parameter is constrained to be $\\Omega_{\\rm DE}^{(0)} = 0.686 \\pm 0.020$ (68\\,\\%\\,CL) \\cite{Planck}. Varying $\\Omega_{\\rm DE}^{(0)}$ in this range, the allowed regions tend to be a little bit narrower in all three models. However, this does not give rise to any qualitatively different change to our results.  We have thus shown that the recent observations of redshift-space distortions are already quite powerful to constrain a number of modified gravity models.  In upcoming future we should obtain more $f \\sigma_8$ measurements that are more precise and/or at different redshifts by larger data of ongoing surveys or near-future projects such as Subaru/FMOS, Subaru/PFS, HETDEX and so on. We hope that we will be able to approach the origin of dark energy in the foreseeable future."
    },
    "1208/1208.4224_arXiv.txt": {
        "abstract": "We use smoothed particle hydrodynamics simulations of massive protostellar discs to investigate the predicted broadening of molecular lines from discs in which self-gravity is the dominant source of angular momentum transport. The simulations include radiative transfer, and span a range of disc-to-star mass ratios between $M_d / M_* = 0.25$ and $M_d / M_* = 1.5$. Subtracting off the mean azimuthal flow velocity, we compute the distribution of the in-plane and perpendicular peculiar velocity due to large scale structure and turbulence induced by self-gravity. For the lower mass discs, we show that the characteristic peculiar velocities scale with the square root of the effective turbulent viscosity parameter, $\\alpha$, as expected from local $\\alpha$-disc theory.  The derived velocities are anisotropic, with substantially larger in-plane than perpendicular values. As the disc mass is increased, the validity of the $\\alpha$ approximation breaks down, and this is accompanied by anomalously large in-plane broadening. There is also a high variance due to the importance of low-$m$ spiral modes.  For low-mass discs, the magnitude of in-plane broadening is, to leading order, equal to the predictions from $\\alpha$ disc theory and cannot constrain the source of turbulence. However, combining our results with prior evaluations of turbulent broadening expected in discs where the magnetorotational instability (MRI) is active, we argue that self-gravity may be distinguishable from the MRI in these systems if it is possible to measure the anisotropy of the peculiar velocity field with disc inclination. Furthermore, for large mass discs, the dominant contribution of large-scale modes is a distinguishing characteristic of self-gravitating turbulence versus MRI driven turbulence. ",
        "introduction": "\\noindent Protostellar discs are at the heart of both pre-main sequence evolution and planet formation theory.  The accretion of the disc material onto young stellar objects (YSOs) is crucial to explaining their phenomenology.  This requires angular momentum to be transported outwards, so that mass may be transported inwards (see e.g. \\citealt{Pringle1981a,Lodato2008,Armitage2011}). Quantitative models of this transport process are needed to describe the evolution of the disc's surface density profile, the protostar's accretion rate, the initial conditions for planet formation \\citep{Chiang2010}, and the subsequent evolution of the planets in the disc. In the simplest description, angular momentum transport is assumed to result from turbulence within the disc, leading to an effective turbulent viscosity \\citep{Shakura_Sunyaev_73},  \\begin{equation} \\nu_{\\rm eff} = \\alpha c_s H. \\label{eq:nu2alpha} \\end{equation}  \\noindent Here, $c_s$ is the sound speed, $H$ is the scale height, and $\\alpha$ is a parameter that characterizes the efficiency of the transport. Constraints from observed disc lifetimes and accretion rates \\citep{Hartmann1998} are consistent with models in which $\\alpha \\sim 10^{-2}$.  Prostostellar disc angular momentum transport is thought to result from two different mechanisms (depending on the location and time in the disc's evolution). In regions of sufficiently high levels of ionization, the magnetorotational instability (MRI) is a plausible source of turbulence \\citep{Balbus1991,Balbus1998,Papaloizou2003}. Further out in the disc (or early on when the disc is quite massive), the disc may be cool and massive enough that the disc's own self-gravity is the dominant source of turbulence \\citep{Lin1987,Laughlin1994}.  This requires that the Toomre $Q$ parameter \\citep{Toomre_1964, Durisen_review},  \\begin{equation} Q = \\frac{c_s \\kappa}{\\pi G \\Sigma} < 1.5-1.7 \\label{toomre_Q} \\end{equation}  \\noindent where $\\kappa$ is the epicyclic frequency ($\\kappa=\\Omega$ for Keplerian discs), and $\\Sigma$ is the disc surface density.  Discs with large $Q$ will undergo radiative cooling towards lower values. The onset of non-axisymmetric instability produces spiral structures in the disc, heating it through shocks and viscous-scale dissipation. Typically, self-gravitating discs will reach a self-regulated, quasi-steady state referred to as marginal instability \\citep{Paczynski1978}, where the heating and cooling are in approximate balance, and $Q\\sim 2$.  This produces a self-sustaining \\emph{gravito-turbulence}, which can be described by an effective $\\alpha$ \\citep{Gammie,Lodato_and_Rice_04}, although there are circumstances where this approximation fails \\citep{Forgan2011}. It is also possible that both MRI and self-gravitating turbulence could be generated at the same time, particularly in the case of layered accretion discs \\citep{Gammie1996,Armitage_et_al_01, Zhu2009,Martin2011}.  Determining {\\em observationally} which mechanism is driving transport at a particular radius in a disc is difficult. Although the condition for the onset of self-gravity (Equation~\\ref{toomre_Q}) is simple and well-understood, measuring $Q$ requires an absolute determination of the gas surface density as a function of radius. This is hard. Most disc mass estimates are derived from dust tracers, which may be systematically biased \\citep{Andrews2007a}. Independent observational discriminants of turbulence would therefore be very valuable.  One promising route is to detect the turbulent broadening of molecular lines in the infrared and/or sub-mm. For local fluid turbulence, elementary arguments \\citep{Balbus1998} suggest that the characteristic velocity perturbations will be of order $v \\sim \\sqrt{\\alpha} c_s \\sim 0.1 c_s$. Using shearing box simulations of MHD turbulence, \\cite{Simon2011} show that this is indeed the approximate value of velocity perturbations at the disc midplane, but that a few scale heights above the midplane they can be as high as $0.5 c_s$, and that a small fraction of the broadening is in fact supersonic.  While extracting turbulent broadening for molecular species directly from the simulations would be ideal, doing so requires an understanding of the disc's density and temperature fields and the abundance of the molecular lines in question.  Modelling these properties in observations is non-trivial, as is the subsequent modelling of radiation transport through the disc.  Projection and resolution effects must also be considered.  Having said this, recent observations have successfully constrained turbulent broadening.  For example, using the CO (3-2) transition with the Sub Millimetre Array (SMA), \\citet{Hughes2011} placed an upper limit of $v < 0.1 c_s$ in the TW Hydra system and determined a broadening value of $v \\approx 0.4 c_s$ in HD 163296.  With the advent of more sensitive millimetre instruments such as ALMA, these constraints will be significantly improved upon.  In this paper, we present preliminary calculations of turbulent velocities in global, three-dimensional smoothed particle hydrodynamics simulations of self-gravitating protostellar discs with radiative transfer \\citep{intro_hybrid}. Our goal is to characterise the turbulent velocity distribution to first-order, and identify any qualitative features in the turbulent velocity field that might enable the discrimination between alternate mechanisms for driving turbulence (although we leave the extraction of synthetic line observations for future work). Existing studies of the MRI and of self-gravitating turbulence suggest that significant differences may exist. Although large-scale magnetic fields play a role in MHD turbulence \\citep{Simon2012}, the MRI can be described to a first approximation as a local angular momentum transport mechanism.  The same is not true of self-gravity in the limit where the disc is massive \\citep{Balbus1999,Lodato2005,Forgan2011}. With this in mind, we seek to identify phenomenological differences which would allow us to distinguish whether line-broadening is MHD-driven or gravity-driven. Such a discriminant would be complementary to estimates based upon derived disc masses, or upon direct imaging searches for spiral structure \\citep{Cossins2010a}. ",
        "conclusions": "\\label{sec:Discussion}   \\noindent We have presented preliminary studies of turbulent line broadening due to self-gravity using global smoothed particle hydrodynamics (SPH) simulations of protostellar discs.  With radiative transfer included in these calculations, we carried out raytracing techniques to measure the local velocity components in the disc plane and transverse to the disc plane (relative to the local sound speed). From these raytracing measurements, we construct a linewidth probability distribution (LPD) which gives the probability of detecting lines broadened to a given width.  These simulations cannot resolve turbulence at scales below the smoothing length, and artificial viscosity effects will also smooth out turbulence to a lesser extent.  That being said, we have detected some promising diagnostics that can be compared to current MRI-driven turbulence simulations but that will require follow-up with high resolution studies.  First, we have found that the LPD is noticeably mass dependent, with the maximum value of $v/c_s$ increasing towards $v/c_s = 1$, where there is a strong cutoff.  Non-local angular momentum transport is evident from strong peaks in the distribution in $v_p$ being caused by low-$m$ spiral structure impinging on the disc radii being observed. This behaviour becomes apparent when the same disc is compared at various radii of observation.  Low mass discs which possess weaker, high $m$ spiral structure retain similar LPDs at all radii, with a systematic separation between the peak of the $v_z$ distribution and the peak of the $v_p$ distribution; $v_p$ peaks near $\\sim 0.04 c_s$, whereas $v_z$ peaks near $\\sim 0.005-0.01 c_s$.  The results from this work can be compared with the work of \\cite{Simon2011} as a way to extract potential differences between the LPDs of MRI-driven turbulence and motions due to self-gravity.  First, comparing the actual mid-plane values of $v/c_s$, we find roughly the same (though somewhat lower) peak $v_p/c_s$ in our low mass discs compared to MRI-driven discs in the ideal MHD limit \\citep{Simon2011}. However, the peak values of $v_p/c_s$ in this work are slightly larger than the dead-zone values of the 4 au simulation of \\cite{Simon2011}.  The largest difference between the LPDs of low mass self-gravitating discs and MRI-driven turbulent discs is the significant separation of $v_z$ and $v_p$ in the self-gravitating case; this robust result is not present in the MRI-driven simulations as $v_z$ and $v_p$ have roughly the same LPD \\citep{Simon2011}.  As such, it would appear to be a good diagnostic for distinguishing between these two angular momentum transport mechanisms in the limit of low disc mass.  In the high disc mass limit, the LPD fluctuates strongly in $v_p$ as the observation radius is changed because spiral arms move in and out of the annulus.  No such fluctuation is detected in $v_z$, but this is possibly due to resolution effects smoothing them out.  The high time variability exhibited by these arms is also apparent in the LPDs as measured across time.  This is another feature that was not observed in the simulations of \\cite{Simon2011}.  The angular momentum transport in these simulations was highly local and there were no low $m$ features observed in the LPDs.  However, these particular simulations were carried out in the local limit; going to larger scales by performing global MRI simulations may introduce the presence of low $m$ features into the LPD \\citep{Beckwith2011, Simon2012}.  In conclusion, the global simulations carried out in this work suggest that one could potentially distinguish between MRI-driven turbulence and self-gravitational effects by observing the variation of turbulent line width with radius (to potentially probe low $m$ structures) and inclination angle.  We must reiterate, however, that our simulations are quite low in resolution, and as such, important phenomenology may not be exhibited.  In addition to this, stellar irradiation - an important driver of line emission from the disc surface - is not included as a radiation source.  We recommend that local simulations be carried out in much the same manner as \\citet{Simon2011}.  These simulations would be capable of confirming the features of gravito-turbulence seen in the low-mass discs, whose angular momentum transport is locally determined \\citep{Forgan2011}.  This includes the systematic separation of $v_z$ and $v_p$.  Also, the increase in resolution afforded by local simulations would allow the study of gravito-turbulent broadening as a function of disc altitude, giving a further means of comparison to MRI turbulence. Stellar irradiation must also be considered in future simulations used for linewidth studies.Finally, it is clear that more sophisticated attempts to make synthetic observations of the linewidth broadening are required to determine the detectability of these features using current and future telescopes.  It is also worth noting that while theory indicates that self-gravitating protostellar discs should exist only in the early embedded phase of protostar formation, this does not present an insurmountable obstacle to observing these features.  Judicious selection of a line with high critical density for activation should allow the midplane to be probed without contamination from the envelope."
    },
    "1208/1208.4633.txt": {
        "abstract": "% {The cooling histories of individual meteorites can be empirically reconstructed by using ages obtained from different radioisotopic chronometers having distinct closure temperatures. For a given group of meteorites derived from a single parent body such data permit the detailed reconstruction of the cooling history of that body. Particularly suited for this purpose are H chondrites because (i) all of them are thought to derive from a single parent body (possibly asteroid (6) Hebe) and (ii) for several specimens precise radiometric ages over a wide range of closure temperatures are available. } % {A thermal evolution model for the H chondrite parent body is constructed by using the cooling histories of all H chondrites for which at least three different precise radiometric ages are available. The thermal model thus obtained is then used to constrain some important basic properties of the H chondrite parent body. } % {Thermal evolution models are calculated using our previously developed code (Henke et al. 2012), which incorporates the effects of sintering and uses new thermal conductivity data for porous materials (Krause et al. 2011). Several key parameters determining the thermal evolution of the H chondrite parent body are varied together with the unknown original location of the H chondrites within their parent body until an optimal fit between the radiometric age data and the properties of the model is obtained. The fit is performed in an automated way based on an `evolution algorithm'. Empirical data for the cooling history of H chondrites are taken from the literature and the thermal model is optimized for eight samples for which radiometric ages are available for at least three different closure temperatures. } % {A set of parameters for the H chondrite parent body is found that yields excellent agreement (within error bounds) between the thermal evolution model and empirical data for the cooling histories of six of the examined eight H chondrites. For two of the samples significant discrepancies exist between model and empirical data, most likely reflecting inconsistencies in the empirical cooling data. The new thermal model constrains the radius and formation time of the H chondrite parent body, and the initial burial depths of the individual H chondrites. In addition, the model provides an estimate for the average surface temperature of the body, the average initial porosity of the material the body accreted from, and the initial $^{60}$Fe content of the H chondrite parent body. } {} ",
        "introduction": "Radiometric ages for chondritic meteorites and their components provide information on the accretion timescale of chondrite parent bodies, and on cooling paths within certain areas of these bodies. However, to utilize this age information for constraining the internal structure, and the accretion and cooling history of the chondrite parent bodies, the empirical cooling paths obtained by dating chondrites must be combined with theoretical models of the thermal evolution of planetesimals. Important parameters in such thermal models include the initial abundances of heat-producing short-lived radionuclides ($^{26\\!}$Al and $^{60}$Fe), which are determined by the accretion timescale, and the terminal size, chemical composition and physical properties of the chondritic planetesimals. Here we use our newly developed code for modelling the structure and thermal evolution of planetesimals of $\\approx100$\\,km size \\citep[][% henceforth called paper I]{Hen11} to constrain the properties and thermal evolution of the H chondrite parent body, and to evaluate the most important parameters that determined its thermal evolution. This is done by finding the optimal fit between a theoretical thermal model and empirical data for the cooling histories of H chondrites. The H chondrite parent body was chosen here as a test case, because it is the only body for which a comprehensive set of thermochronological data is available. For eight H chondrites in particular radiometric ages exist that were obtained using at least three distinct radioisotope chronometers characterised by different closure temperatures. The cooling history of these samples, therefore, is well constrained.  %--------------------------------- \\begin{table*}  \\caption{Key data for cooling history of selected H chondrites}  \\begin{tabular}{@{}llllllllll@{}} \\hline\\hline \\noalign{\\smallskip} Meteorite & Estacado & Guare\\~na & Kernouv\\'e & Richard- & Allegan & Nadiabondi & Forest Vale & Ste. Mar- \\\\ &          &         &            & ton      &         &            &             & guerite \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} type    & H6 & H6 & H6 & H5 & H5 & H5 & H4 & H4 & \\\\ \\noalign{\\smallskip} & \\multicolumn{7}{c}{Hf-W$^{\\rm f}$ (metal-silicate)} \\\\ \\noalign{\\smallskip} radiometric age & $4557.3\\pm 1.6 $ & & 4557.9$\\pm$1.0 & 4561.6$\\pm$0.8 & & & & 4565.6$\\pm$0.6& Ma\\\\ temperature     & 1150$\\pm$75 & & 1150$\\pm$75 & 1100$\\pm$75 & & & & 1070$\\pm50$ &  K \\\\ % \\noalign{\\smallskip} & \\multicolumn{7}{c}{Pb-Pb$^{\\rm g}$ (pyroxene-olivine)} \\\\ radiometric age &  &  &  4537.0$\\pm$1.1 &  &  & 4558.9$\\pm$2.3 &  & 4564.3$\\pm$0.8 & Ma\\\\ temperature     &  &  &  1050$\\pm100$    &  &  & 1050$\\pm100$    &  & 1050$\\pm100$    &  K \\\\ % \\noalign{\\smallskip} & \\multicolumn{7}{c}{U-Pb-Pb$^{\\rm e}$ (phosphates)} \\\\ \\noalign{\\smallskip} radiometric age & 4559$\\pm$2 & 4504.4$\\pm$0.5 & 4522.5$\\pm$2.0 & 4551.4$\\pm$0.6 & 4550.2$\\pm$0.7 & 4555.6$\\pm$3.4 & 4560.9$\\pm$0.7 & 4562.7$\\pm$0.6 &  Ma\\\\ temperature$^{\\rm a}$     & 720$\\pm$50 & 720$\\pm$50 & 720$\\pm$50 & 720$\\pm$50 & 720$\\pm$50 & 720$\\pm$50 & 720$\\pm$50 & 720$\\pm$50 &  K \\\\ \\noalign{\\smallskip} & \\multicolumn{7}{c}{Ar-Ar$^{\\rm e,h}$ (feldspar)} \\\\ \\noalign{\\smallskip} radiometric age & 4465$\\pm$5 & 4458$\\pm$10 & 4499$\\pm$6 & 4525$\\pm$11 & 4541$\\pm$11 & 4535$\\pm$10 & 4552$\\pm$8 & 4562$\\pm$16 &  Ma\\\\ temperature$^{\\rm b}$     & 550$\\pm$20 & 550$\\pm$20 & 550$\\pm$20 & 550$\\pm$20 & 550$\\pm$20 & 550$\\pm$20 & 550$\\pm$20 & 550$\\pm$20 &  K \\\\ \\noalign{\\smallskip} & \\multicolumn{7}{c}{Pu-fission tracks (pyroxene)} \\\\ \\noalign{\\smallskip} calculated age$^{\\rm c}$ & 4465$\\pm$5 & 4458$\\pm$10 & 4499$\\pm$6 & 4525$\\pm$11 & 4541$\\pm$11 & 4535$\\pm$10 & 4552$\\pm$8 & 4562$\\pm$16 &  Ma \\\\ temperature & $550\\pm25$ & $550\\pm25$ & $550\\pm25$ & $550\\pm25$ & $550\\pm25$ & $550\\pm25$ & $550\\pm25$ & $550\\pm25$ & K \\\\ & \\multicolumn{7}{c}{Pu-fission tracks (merrilite)} \\\\ \\noalign{\\smallskip} calculated age$^{\\rm c}$ & 4401$\\pm$10 & 4402$\\pm$14 & 4438$\\pm$10 & 4469$\\pm$14 & 4490$\\pm$14 & 4543$\\pm$20 & 4544$\\pm$14 & 4550$\\pm$17 &  Ma \\\\ temperature & $390\\pm25$ & $390\\pm25$ & $390\\pm25$ & $390\\pm25$ & $390\\pm25$ & $390\\pm25$ & $390\\pm25$ & $390\\pm25$ & K \\\\ \\noalign{\\smallskip} & \\multicolumn{7}{c}{Pu-fission tracks (pyroxene --- merrillite)} \\\\ \\noalign{\\smallskip} time interval$^{\\rm d,e}$ & 64$\\pm$9 & 56$\\pm$9 & 61$\\pm$8 & 56$\\pm$9 & 51$\\pm$9 & $>-8$  & $<8$ & 12$\\pm$11 &  Ma \\\\ cooling rate & 2.5$\\pm$0.4 & 2.8$\\pm$0.6 & 2.6$\\pm$0.5 & 2.9$\\pm$0.5 & 3.2$\\pm$0.6 & (fast) & $>20$ & 12.8$\\pm$6 &  K/Ma \\\\ %\\noalign{\\smallskip} %         & \\multicolumn{7}{c}{metallographic } \\\\ %\\noalign{\\smallskip} % cooling rate$^{\\rm d,e}$    & 10 & 4 & 10 & 20 & 15 & 50 & 1000 & 1000 %    &  K/Ma \\\\ \\noalign{\\smallskip} \\hline \\end{tabular}  \\medskip{\\scriptsize Notes: \\\\ $^{\\rm a}$~\\begin{minipage}[t]{.975\\hsize} see \\citet{Cer91} \\end{minipage}  $^{\\rm b}$~\\begin{minipage}[t]{.975\\hsize} \\citet{Tri03} \\end{minipage}  $^{\\rm c}$~\\begin{minipage}[t]{.975\\hsize} Calculated age at 390\\,K from time interval between Pu-fission track (merillite, 390\\,K) and Pu-fission tracks at 550\\,K (pyroxene) compared to Ar-Ar feldspar age at 550\\,K \\end{minipage}  $^{\\rm d}$~\\begin{minipage}[t]{.975\\hsize} Time-interval for Pu-fission track cooling rate from 550--390\\,K, metallographic cooling rate 800--600\\,K \\end{minipage}  $^{\\rm e}$~\\begin{minipage}[t]{.975\\hsize} Data from U-Pb-Pb: \\citep{Goe94}, Ar-Ar: \\citet{Tri03}, metallographic cooling rates: \\citet{Tay87} \\end{minipage}  $^{\\rm f}$~\\begin{minipage}[t]{.975\\hsize} Hf-W ages are from \\citet{Kle08} and were re-calculated relative to the $^{182}$Hf/$^{180}$Hf of the angrite D'Orbigny, which has a Pb-Pb age of $t =4563.4\\pm0.3$\\,Ma \\citep{Kle12}. \\end{minipage}  $^{\\rm g}$~\\begin{minipage}[t]{.975\\hsize} \\citet{Bou07} \\end{minipage}  $^{\\rm h}$~\\begin{minipage}[t]{.975\\hsize} Recalculated for miscalibration of K decay constant \\citep[explanation see text and][]{Tri03} \\end{minipage}  }  \\label{TabDatChondCool} \\end{table*} %--------------------------------  The well preserved and smooth cooling histories of many H chondrites were used to argue that the H chondrite parent body has not been subjected to catastrophic collisions during the first $\\approx100$\\,Ma of its history. In that case an onion shell structure develops in which the degree of thermal metamorphism is a function of depth within the parent body; material of petrologic type H6 is found in the central part, while type H3 material is located at the surface. Based on the assumption of an onion shell structure, a number of thermal evolution models were constructed that used increasingly more complex physical input \\citep{Min79,Miy81,Ben96, Akr98,Tri03,Hev06,Sah07,Kle08,Har10}. All the previous thermal models of the H chondrite parent body found a radius of about 100\\,km and an accretion time of about 2\\,Ma after CAI formation. The success of these models in reconstruction the empirical cooling histories of the H chondrites gave much support to the hypothesis that catastrophic collisions did not significantly disturb the initial thermal structure of the H chondrite parent body. It should be noted, however, that some hints exist for the occurrence of impacts \\citep[see][and references therein]{Wit10}, but these do not appear to have resulted in catastrophic disruption.  These previous thermal models, despite their success in reconstructing the observed empirical cooling rates for several H chondrites, fall short in providing insights into the physical processes occurring during the thermal history of a $\\approx100$\\,km-sized planetesimals. This motivated us to use our new thermal evolution model (described in paper I) to find the optimal fit between the empirical cooling history of the H chondrites and the thermal model. Our new thermal model applies recent data on the thermal conductivity of granular material \\citep{Kra11,Kra11b} and on the compaction by cold isostatic pressing before the onset of sintering \\citep{Gue09}. The modelling of the sintering process is based on the same kind of theory \\citep{Rao72} as in \\citet{Yom84}, but with some improvements. This theory does not consider more recent approaches for modelling hot pressing for technical processes \\citep[e.g.][]{Arz83,Fis83,Lar96,Sto99} but appears more appropriate for the lower pressure regime relevant for asteroids. The treatment of heat conduction of the chondritic material and its sintering under pressure and at high temperature still has to be considered as rather crude, but at least is based on physical concepts that are successfully applied in other fields and in some laboratory measurements.  We do not consider a possible melting of the body, because the H chondrite parent body is thought to have remained undifferentiated. Our results will show, however, that this assumption may not be entirely valid, because the calculated central temperature of the H chondrite parent body is only slightly below the solidus temperature of chondritic material \\citep[1200\\,K, ][]{Fei97}.  Our model considers more properties of the asteroidal body than traditional approaches, and so fitting of the model to the observed cooling histories is more difficult because of the much larger parameter space. Therefore we develop an automated method to perform such a fit that is based on the \u0093evolution algorithm\u0094 that mimics the concepts that are thought to rule biological evolution in nature. This is some kind of \u0092intelligent\u0092 trial-and-error method by which the optimum of some quality function can be found. Its advantage is that it does not require any kind of good behaviour of the quality function. It will be shown that this method \\citep[in the version described in][]{Cha95} can successfully be applied to our problem.  For fitting the model we exclusively use meteorites for which at least three age determinations at different closure temperatures are available. These meteorites do not only constrain the slope but also the curvature of an individual cooling path at a certain location within the parent body. Within the framework of the analytic model of \\cite{Miy81}, such a data set for a single meteorite would already completely determine the radius and the formation time of the body. However, in view of the presumably complex thermal evolution of meteorite parent bodies, a larger dataset for several meteorites is necessary for a reliable fit. For the H chondrites, 31 data from 8 meteorites (2 meteorites with 5 data, 3 with 4 data, 3 with 3 data) are available for this purpose. This allows determining all the important parameters of the H chondrite parent body by an optimization procedure.  The plan of this paper is as follows: In Sect.~\\ref{SectCoolHist} we describe the available empirical data. In Sect.~\\ref{SectEvoMod} we give a brief overview how we calculate thermal evolution models. Section~\\ref{SetFitProc} describes the optimisation method.  Section \\ref{SectFitHChond} gives our results and conclusions. The paper closes with some final remarks in Sect.~\\ref{SectConclu}.    %******************************************************************************* ",
        "conclusions": "\\label{SectConclu}  We combined the code for modelling the internal structure and evolution of planetesimals of the 100\\,km size class \\citep{Hen11} with a code for optimisation of a quality function that is based on an evolution algorithm. This allows finding a consistent set of parameters that describe the properties and evolution of a planetesimal and the burial depths of the meteorites. A such the thermal evolution of the model body reproduces very closely the empirically determined cooling history of all H chondrites with sufficient available data.  The results show, that the ``onion shell'' model successfully reproduces the thermal evolution history of all the meteorites used in this study. For six of the eight meteorites a burial depth can be found for which  the empirical cooling data can be fitted in an excellent way with the local evolution of temperature at that burial depth. For the remaining two cases it is still possible to find a burial depth that results in a good fit if the Pu fission-track datum of Nadiabondi is rejected. This confirms earlier conclusions that the thermal evolution of the H chondrite parent body and its initial thermal structure has not been disturbed by catastrophic impacts.  The reconstructed properties of the parent body of the H chondrites are comparable to what has been found in other model calculations. The different published models treated the problem in different kinds of approximations, solving the heat conduction equation being almost the only feature common to all models, and used different values for the material properties. Our estimated size of 123\\,km of the parent body is somewhat bigger compared to the 100\\,km found in other models, and requires a slightly earlier formation time of 1.88\\,Ma compared to the 2.2 -- 2.5\\,Ma found in other models. Such minor differences are certainly due to the different approaches used.  Our study shows, that the input physics of the models has to be improved, before more reliable data on the parent body can be obtained. Such improvements are currently underway in our group. In particular the heat conductivity has to be calculated from the specific properties of chondritic material. It will also be necessary to obtain more precise radiometric ages for a wider range of samples.  The advantage of our method is that it allows for a simultaneous and consistent fit of several model parameters to a big number of observational restrictions and is very flexible if the kind of data to be included in the optimization process is to be changed. A minor shortcoming of the optimisation method used is that a considerable number of complete evolution models has to be calculated during the course of the optimisation process, but our implicit solution method is fast enough that no restriction as to the applicability of the method to our problem results from this.   %------------------------------------------------------------------------------"
    },
    "1208/1208.2717_arXiv.txt": {
        "abstract": "Deep HST broad-band images taken with ACS and WFPC2 of the giant ($\\sim 1000$~AU diameter) dark silhouette proplyd 114-426 in the Orion Nebula  show that this system is tilted, asymmetric, warped and photoevaporated. The exquisite angular resolution of ACS allows us to map the distribution of dust grains at the northern translucent edge of the disk, dominated by the photoevaporative flow. Using the Mie theory for standard circumstellar disk grains, we find evidence for a spatial gradient in grain size. The typical dust radius, $\\simeq 0.2-0.7~\\mu$m (less than what reported by previous studies) becomes smaller as the distance from the disk center increases, consistent with the expectations for the dynamic of dust entrained in a gaseous photoevaporative wind. Our analysis of the disk morphology and location within the nebula indicates that this system is photoevaporated by the diffuse radiation field of the Orion Nebula, while being shielded from the radiation coming directly from the central Trapezium stars. We  estimate the mass-loss rate from the disk surface and the time-scale for total disk dissipation, which turns out to be of the order of $10^4$yr. Such a short time, of the order of 1/100 of the cluster age, indicates that this system is seen on the verge of destruction. This is compatible with the exceptional nature of the disk, namely its combination of huge size and low mass. Finally, we briefly discuss the viability of possible mechanisms that may lead to the peculiar morphology of this system: external UV flux, binary star and past close encounter. ",
        "introduction": "Early \\emph{Hubble Space Telescope} (HST)  images of circumstellar disks in the Orion Nebula Cluster have provided the first direct evidence that protoplanetary disks in rich young stellar clusters may be photo-evaporated by UV radiation of nearby massive stars  \\citep{ODellWen94}. The typical signature of photo-evaporation is a bright ionized cusp pointing to the brightest stars of the cluster, $\\theta^1$Ori-C in particular. The cusps always surround a central low-mass star, that may occasionally be encircled by a compact dark disk seen in silhouette against the nebular background, or against the cusp itself. The most recent atlas of protoplanetary disks (``proplyds'') in the Orion Nebula \\citep{Ricci} counts 178 of these comet-shaped sources. The same atlas also counts other 36 disks which do not show evidence of photoevaporation, being seen as pure dark silhouettes either in absorption against the nebular background (28 sources) or as dark lanes at the midplane of some diffuse bipolar emission (8 sources). Disk photo-evaporation naturally raises the issue of disk survival in the harsh environment of young massive clusters  \\citep[see e.g.][]{ODell+08}. The frequency and spatial distribution of disks with and without evidence of photo-evaporation within the same cluster provides an excellent  opportunity to investigate how environmental factors may affect the path to planet formation. \\\\   The basic physical process at the origin of disk photo-evaporation in massive clusters has been clarified by  \\citet{Johnstone}. The FUV (6 eV $\\le h\\nu \\le$ 13.6 eV; 2000\\AA$\\ge\\lambda\\ge$ 912\\AA) radiation reaches nearly undisturbed the disk surface, creating a photodissociation region (PDR) with temperature $T\\sim100-1000$\\,K. Depending on the radial distance from the ionizing stars, the heated disk surface may reach thermal velocity higher than the escape velocity. The resulting neutral outflow, largely composed by molecular $H_2$ gas dragging a non negligible amount of small dust grains \\citep{Owen} soon encounters the EUV  ($h \\nu \\ge$ 13.6 eV; $\\lambda \\le$ 912\\AA) radiation coming directly from the ionizing stars, creating a ionization front at $T\\simeq10^4$~K, i.e. the bright ionized cusp. This model can explain both the general morphology of photoevaporated disks as well as more peculiar features, like the presence of [OI] emission from photodissociation of OH radicals  produced by $H_2$ heating and dissociation observed in one of the most prominent objects, 182-413 \\footnote{In this paper we shall adopt the coordinate-based naming convention of \\citet{ODell+08}.} \\citep{StorzerHollenbach99, Bally+00}.  It is generally assumed that disk photo-evaporation in rich clusters, as revealed by the presence of photo-ionized cusps, can occur only if the disk is directly exposed to the EUV flux from the massive stars, consistently with the fact that the majority of photo-evaporated disks in the Trapezium cluster lies in the immediate surroundings of $\\theta^1$Ori-C. Dark silhouette disks, lacking any evidence of photo-ionization, are believed to lie within the foreground Veil of the Orion Nebula, shielded by EUV radiation \\citep{McCaughreanOdell96}. However, dark silhouette disks can also be photo-evaporated by FUV radiation \\citep{Adams+04} although the products of photo-dissociation in these systems have never been seen \\citep{ODell+08}. This because the parameter controlling photo-evaporation is the intensity of the FUV flux, which can be expressed in units of \\emph{Habing Flux} (the estimate average flux in the local interstellar medium), defined as $G_0 = 1.6 \\times 10^{-3} $erg cm$^{-2}$ s$^{-1}$. In the inner core of the Trapezium, the FUV flux is elevated, reaching $G = 10^5 G_0$ at the Orion Bar \\citep{Robberto+02}, about  110\\arcsec or 0.22~pc south of $\\theta^1$Ori-C \\citep[assuming a distance $d = 414$~pc][]{Menten+07}\\footnote{while \\citet{Menten+07} provide the most precise estimate, \\citet{ODellHenney08} consider a distance $d=430\\pm20$~pc to be more accurate. Our estimates  do not depend on the exact value of the assumed distance.}. Sources in the outskirt of the region, shielded from the EUV photons emitted e.g. by $\\theta^1$Ori-C, may still receive scattered EUV and, especially, FUV radiation from the nebular environment.  The most spectacular case of dark silhouette in Orion is the well know 114-426 disk \\citep{McCaughrean98}. With an apparent diameter of over $2\\arcsec$, which makes it visible also from the ground, this disk spans about $\\sim 1000$ AU, 10 times the size of our Solar System or of other typical photo-ionized disks in Orion. It is seen nearly edge-on, the presence of the central star being noticeable through scattered polar emission, and is considered to be in a relatively quiescent, non photo-evaporated phase \\citep{Throop+01}.  \\citet{McCaughrean98} estimate for the central star an initrinisic brighness $K\\sim9.5^m$, corresponding to a 1.5M$_\\odot$ star at the distance and age of the Orion cluster, but the uncertainites are large. Using HST/NICMOS high-resolution near-infrared images of 114-426, \\cite{McCaughrean98} find that the major axis of the disk in the Pa$\\alpha$ at 1.87 $\\mu$m line is $\\sim$20\\,$\\%$ smaller than at 0.6\\,$\\mu$m. From this, they deduce a minimum disk mass of 10\\,M$_\\oplus$, and possibly $\\geq 5 \\times 10^{-4}$\\,M$_{\\odot}$. On the other hand,  114-426 has not been detected by the Submillimeter Array survey  \\citep{Mann} at either 880\\,$\\mu$m or 1330\\,$\\mu$m. This non-detection places an upper mass limit of $ 1.2 \\times 10^{-2} M_{\\odot}$ (0.5 - 12 $M_{Jup}$), confining the disk mass into a range of two orders of magnitude.  This object has attracted interest because the variation of the disk size with wavelength has allowed to probe the properties of the disk grains, exploiting the wavelength dependence of the extinction of the bright nebular background through the translucent disk edges \\citep{Throop+01}. If the grain sizes were much smaller than the observed wavelengths, as in the Rayleigh scattering limit, shorter wavelengths would be scattered much more efficiently than  longer ones: in this case the disk would appear smaller at long wavelengths. Viceversa, if  the grain sizes are of the same order of the wavelength of observation, the disk size is expected to remain nearly constant with wavelength (``gray'' scatter).  Using observations of  114-426 in the H$\\alpha$ (0.656 $\\mu$m; HST/WFPC2), Pa$\\alpha$ (1.87 $\\mu$m; HST/NICMOS1) and Br$\\alpha$ (4.05 $\\mu$m; Keck/NIRSPEC), \\citet{Shuping} found evidence for chromatic extinction at the northern translucent edge of the disk, deriving a typical grains diameter $\\gtrsim 1.9 \\mu$m, but not $\\gg$ 4 $\\mu$m.  This represents one of the first evidences for grain growth in protoplanetary disks,  the initial step in the process of planet formation \\citep{Natta+07}.  In this paper we use the most recent HST/ACS images of the 114-426, together with a HST/WFPC2 image taken in the U-band, to clarify the disk structure and its physical status. Contrary to what is commonly assumed, we show that the disk appears strongly affected by the action of the environment and is photoevaporated. Its unusually large size and morphology (warped and asymmetric with respect to the central star) may result from photoevaporation. The translucent northern edge of the disk appears originated by the outflow of neutral material rather than the truncated edge of an undisturbed disk. We map the variation of extinction in the translucent region, showing how the typical grain size variation is compatible with an evaporative flow. This allows us to estimate the mass-loss rate from the disk surface and the disk lifetime. We also show that the disk is enshrouded by a large scale ($\\sim0.03$~pc)  cocoon of dust  with a bright ridge facing the side opposite to the cluster center. This suggests that the disk photoevaporation is driven by the diffuse radiation field rather than by direct irradiation from the massive cluster stars.  In Section~\\ref{sec:Observations} we briefly present our data; in Section~\\ref{sec:Morphology} we illustrate the geometry of the source, moving from the inner structure to the most extended features and the nebular surroundings; in Section~\\ref{sec:Dust}  we analyze the grain size composition and distribution in the photoevaporated flow; in Section~\\ref{sec:Discussion} we derive the main physical parameters of the mass loss and discuss the possible heating sources. Our conclusions are presented in Section~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions} We have used a set of HST images taken with WFPC2 and ACS to describe the complex morphology of the most prominent dark silhouette disk in the Orion Nebula, 114-426. The disk appears tilted, asymmetric and warped, with clear signs of photoevaporation of neutral material. Mapping the distribution of dust grains at the northern translucent edge of the disk using the Mie theory for standard circumstellar disk grains, we have found evidence for a spatial gradient of the grain radius, from 0.2 to 0.7~$\\mu$m, with grains becoming smaller as the distance from the disk increases. These values are slightly below the minimal 1.9~$\\mu$m grain size estimated by Shuping et al. (2003), but still compatible considering the difficulty of observing this source at infrared wavelengths. Overall, the disk morphology and location in the nebula indicate that this system is photoevaporated by the diffuse radiation field of the Orion Nebula, while being shielded from the radiation coming directly from the central Trapezium stars. Using our derived grain properties, we  have estimated the mass loss rate from the disk surface and the time-scale for total disk dissipation, which turns out to be of the order of $10^4$yr. This short time scale, of the order of 1/100 of the cluster age,  seems compatible with the exceptional nature of this system, namely its huge size and low mass. Finally, we briefly discuss the viability of possible mechanisms that may lead to the peculiar morphology of this system: external UV flux, binary star and past close encounter.  \\vskip 0.5cm {\\sl Acknowlegments}: the authors would like to thank S. Lubow, I. Pascucci, J. Pringle and K. Stapelfeld for helpful discussions on this fascinating source, and the referee W. Henney for prompt and useful feedback. Zolt Levay provided support with image production. Support for GO program 10246 was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555."
    },
    "1208/1208.4131_arXiv.txt": {
        "abstract": "At the moment of deepest core collapse, a star cluster core contains less than ten stars.  This small number makes the traditional treatment of hard binary formation, assuming a homogeneous background density, suspect.  In a previous paper, we have found that indeed the conventional wisdom of binary formation, based on three-body encounters, is incorrect.  Here we refine that insight, by further dissecting the subsequent steps leading to hard binary formation. For this purpose, we add some analysis tools in order to make the study less subjective. We find that the conventional treatment does remain valid for direct three-body scattering, but fails for resonant three-body scattering.  Especially democratic resonance scattering, which forms an important part of the analytical theory of three-body binary formation, takes too much space and time to be approximated as being isolated, in the context of a cluster core around core collapse.  We conclude that, while three-body encounters can be analytically approximated as isolated, subsequent strong perturbations typically occur whenever those encounters give rise to democratic resonances.  We present analytical estimates postdicting our numerical results.  If we only had been a bit more clever, we could have predicted this qualitative behaviour. ",
        "introduction": "In our previous paper, \\cite{Tanikawa11}, hereafter referred to as Paper I, we started to investigate in detail the formation mechanism of the first hard binary during core collapse of a dense star cluster. While many studies have appeared that have focused on the macroscopic aspects of core collapse, during the last fifty years, to the best of our knowledge our paper was the first one to address the microscopic aspects, including the actual reaction network of the stellar encounters that gave rise to the formation of a hard binary.  In that study, we encountered two surprising deviations from what had become accepted as the standard picture of binary formation in core collapse.  First, in many cases more than three bodies are directly and simultaneously involved in the production of the first hard binary.  Second, we concluded that the core at deepest collapse was smaller than expected before, typically containing half a dozen stars or less.  In contrast, in the standard picture developed in the nineteen eighties, it was first assumed that the formation of hard binaries was essentially a three-body process, whose rate could be estimated assuming the typical density and velocity dispersion in the core. Second, it was concluded that the core would bounce around the time its membership had dropped to a few dozen stars. In Paper I, we cited papers by \\cite{Goodman84,Goodman87}.  An additional reference is \\cite{Hut85}, where analytical arguments were used to predict that three-body binary formation would reverse core collapse when the core shrank to contain of order 100 stars (80 in their section IVa, and 150 in their section IVbii).  They also quoted simulations by \\cite{McMillan84} which showed core collapse to be reversed when the core contained 25 stars.  The two main flaws in the traditional picture are related.  Given that the fluctuations in thermodynamic properties in a group of only a few stars are far larger than in a group of, say, thirty stars, the concept of a homogeneous temperature (or velocity dispersion) in the core is no longer valid for such a small core. Also, in a core containing only, say, five stars it is quite likely that all five are involved in the formation of a hard binary, with possibly some of the stars just outside the core also making a strong presence felt during a pass through the core.  Encouraged by the fact that the standard story of hard binary formation needed to be corrected on at least these two quite fundamental points, we continued our investigation, focusing in on only one of the many runs reported in Paper I, in an attempt to get further to the bottom of what is actually happening during core collapse in microscopic detail.  Not wanting to introduce any bias, we decided to simply take the very first case described in Paper I.  The central new technique, introduced in Paper I, was to plot all pair-wise distances for all stars in the core, as a function of time, during a short period of time just before hard binary formation. Using this technique, and interpreting the results by eye, was only feasible given the very small number of stars in the core.  Together with visual interactive inspection of the 3-D orbits of the stars in the core, our new technique allowed a determination of roughly how many stars were involved at each time during the successive stages leading to the formation of the first hard binary.  In this paper, we move beyond the detection of the new physics reported in Paper I, i.e. many-body binary formation, in order to perform a more detailed and quantitative analysis of this binary-formation process. In particular, we devote efforts to finding stars involved with the binary formation more objectively, and to reveal what kind of subsystems these stars construct, and how these stars and subsystems interact with each other. For this purpose, we introduce two other new techniques. The first one is the use of work functions, and the second one is a form of subcluster analysis. In addition, we have employed a better interactive visualisation tool, in the form of an {\\tt open-GL} program. These tools will be useful to make clear binary formation in more realistic and complicated $N$-body simulations in which stars have different masses and experience internal evolution.  In the process of applying these new tools, we again found new physics: while the main conclusions of Paper I hold, we now understand in more detail exactly why they hold. The main reason is the presence of democratic resonance interactions, a concept introduced by \\cite{Hut82}, which is a kind of encounter between a hard binary and a single intruder in which a third body is temporarily bound to the binary, in such a way that the subsequent motion cannot be described as a hierarchical triple system. In contrast, we found that the traditional perturbative treatment is in fact satisfactory for direct three-body interactions.  It is only because democratic resonance interactions last long and take up a large fraction of the space in the core that they will typically undergo strong encounters with other stars before a democratic resonance is finished.  This new paper has two main aims: first, to illustrate the new diagnostic tools and their uses (with a long-term goal of making this kind of analysis more streamlined and automatic), and, second, to explain our developed understanding of the new dynamical processes which we are exploring.  This main finding is described in more detail in the discussion and conclusion sections below.  The next two sections focus on a summary of what we found out about the first run in Paper I; and on what we learned in our new analysis in this paper, respectively.  The section after that extends the analysis of Paper I to earlier times, where interesting processes were already happening that had not been flagged in Paper I. The paper finishes with a section of theoretical discussion, and then a summary of our conclusions and some outlook. ",
        "conclusions": "We have dissected the spacetime history of the formation of the first hard binary in a $1024$-body run, in microscopic detail. This paper is the second in a series, following Paper I in which we presented the first such microscopic observation of hard binary formation during core collapse.  The single run that we have investigated here in great detail is the very first run we presented in that paper.  The main improvements over Paper I are: \\begin{enumerate} \\item[1.] We have introduced a new type of reaction diagram, somewhat similar to Feynman diagrams in perturbative quantum field theory calculations, and also similar to what was used by \\citet{1983ApJ...268..319H} in the top part of their fig. 3 (section 3.2). \\item[2.] We have introduced a new tool, in the form of work functions (section 3.1). \\item[3.] We have introduced another new tool, subcluster analysis (section 3.1). \\item[4.] We have highlighted the central role played by democratic resonances, especially three-body resonances whose longevity makes it likely that many-body interactions take place in the core of a star cluster around core bounce (section 3.7 and 5.1). \\item[5.] We have traced the network of reactions leading to the initial formation of the first hard binary back to earlier times, showing a complexity significantly larger even than what we had already unearthed in Paper I (section 4). \\item[6.] We have provided a new qualitative argument to derive the delay of core bounce, compared to standard expectations, based on the delay of hard binary heat production after formation (section 5.4). \\end{enumerate}  It is interesting to note that we have employed four distinct levels of analysis: \\begin{enumerate} \\item[1.] {\\it Visual analysis} of the motions of stars in the core, using an interactive visualisation tool in the form of an {\\tt open-GL} program (we have not stressed this initial phase, but it has helped to guide our intuition and to resolve ambiguities). \\item[2.] {\\it Geometric analysis} based on pairwise distances between interacting stars. \\item[3.] {\\it Energetic analysis} based on the binding energies of pairs and higher-order multiple stars. \\item[4.] {\\it Dynamic analysis} based on energy transfer between tuples of stars. \\end{enumerate} These steps lead to a compressed schematic rendition in the form of the Feynman like diagram depicted in figure \\ref{fig:diagram}.  The next step in our explorations will be to extend the applications of our new techniques to a large number of $N$-body core collapse simulations, for different values of $N$.  In order to do so, much of the analysis presented here will have to be automated.  Ideally, all of the figures presented here would be generated automatically by a single analysis package.  In practice, the development of such a package will remain a formidable challenge for quite a while to come.  A more modest step would be to develop improved tools to help generate many of the figures semi-automatically, requiring far less time and energy than has been the case for Paper I and the current paper, by providing better graphics tools and other diagnostic tools covering the physical properties of the core.  A next step could be to generate a kind of artificially intelligent module that is trying to guess when an interesting network of reactions starts and ends, around the time of core collapse, and whether such a network includes the formation of a surviving hard binary.  Using such a tool would still require human supervision to check whether the results make sense, and to arbitrate in ambiguous situations.  Ideally, after one or more steps, we could then build a software system that fully automatically would produce all the diagrams presented in this paper for any run, including the one introduced here that resembles a Feynman diagram.  The results that we have presented could in principle be obtained from an $N$-body simulation code like {\\tt NBODY6}, which contains modules that allow the user to output logs with information about binaries and their hierarchy \\citep{Aarseth01}. An analysis of these logs are expected to produce the same result as we have obtained (for the same numerical orbits), if the user has a way to deal with the huge amount of data that would be produced. What we add here is a set of tools that enable the user to analyse those kinds of data.  We could take a further step, and add more realistic effects to our simulations, such as primordial binaries. It would be interesting to elucidate whether or not the presence of hard primordial binaries tends to suppress the formation of new binaries."
    },
    "1208/1208.4846_arXiv.txt": {
        "abstract": "We present {\\it Herschel}-PACS spectroscopy of the [OI]63\\,$\\mu$m far-infrared cooling line from a sample of six unlensed and spectroscopically-confirmed 870\\,$\\mu$m-selected submillimetre (submm) galaxies (SMGs) at $1.1<z<1.6$ from the LABOCA Extended Chandra Deep Field South (ECDFS) Submm Survey (LESS). This is the first survey of [OI]63\\,$\\mu$m, one of the main photodissociation region (PDR) cooling lines, in SMGs. New high-resolution ALMA interferometric 870\\,$\\mu$m continuum imaging confirms that these six {\\it Herschel}-targeted SMG counterparts are bona fide sources of submm emission.  We detect [OI]63\\,$\\mu$m in two SMGs with a SNR$\\gsim3$, tentatively detect [OI]63\\,$\\mu$m in one SMG, and constrain the line flux for the non-detections. We also exploit the combination of submm continuum photometry from 250--870\\,$\\mu$m and our new PACS continuum measurements to constrain the far-infrared luminosity, L$_{\\rm FIR}$, in these SMGs to $\\lesssim30$\\,per cent. We find that SMGs do not show a deficit in their [OI]63\\,$\\mu$m--to--far-infrared (FIR) continuum luminosity ratios (with ratios ranging from $\\simeq0.5$--1.5\\,per cent), similar to what was seen previously for the [CII]158\\,$\\mu$m--to--FIR ratios in SMGs.  These observed ratios are about an order of magnitude higher than what is seen typically for local ultra luminous infrared galaxies (ULIRGs), which adds to the growing body of evidence that SMGs are not simply `scaled up' versions of local ULIRGs.  Rather, the PDR line--to--L$_{\\rm FIR}$ ratios suggest that the star formation modes of SMGs are likely more akin to that of local normal (lower-luminosity) star-forming galaxies, with the bulk of the star formation occurring in extended galaxy-scale ($\\sim$kpc) regions.  These observations represent the first step towards a census of the major PDR cooling lines in typical SMGs that will be attainable with ALMA, enabling detailed modelling to probe the global properties of the star formation and the evolutionary status of SMGs. ",
        "introduction": "Far-infrared bright galaxies comprise some of the most luminous galaxies in the local Universe, but they contribute just 0.3 per cent of the total local galaxy luminosity density \\citep{SandMir96}. At first glance it appears that they are an unimportant component of the galactic zoo, however, sensitive (sub)millimetre cameras have revealed that ultraluminous infrared galaxies (ULIRGs; L$_{\\rm FIR}\\geq10^{12}$\\,L$_\\odot$) are much more abundant at high redshifts ($z\\gtrsim2$). Since their discovery (\\citealt{Smail97}; \\citealt{Hughes98}; \\citealt{Barger98}), nearly 13 years of piece-meal survey work has so far produced samples totalling $\\sim500$ 850--1200-selected (sub)mm galaxies (SMGs; e.g.~\\citealt{Coppin06}; \\citealt{Austermann10}; \\citealt{Weiss09}), with extreme star formation rates (SFRs) of $\\sim100$--1000\\,M$_\\odot$\\,yr$^{-1}$ (e.g.~\\citealt{Kovacs06}; \\citealt{Coppin08}; \\citealt{Magnelli12}) and fuelled by vast reservoirs of cold molecular gas (M$_\\mathrm{H_{2}}\\sim5\\times10^{10}$\\,M$_\\odot$; \\citealt{Greve05}; \\citealt{Tacconi06,Tacconi08}; \\citealt{Bothwell12}).  Spectroscopic follow-up of a subset of SMGs suggests that the volume density of ULIRGs increases $\\simeq1000$ times out to $z\\sim2$ (\\citealt{Chapman05}; \\citealt{Banerji11}). Therefore, luminous obscured galaxies likely dominate the total bolometric emission from star formation at early epochs  (e.g.~\\citealt{Pascale09}; \\citealt{Marsden09}; \\citealt{Eales09,Eales10}; \\citealt{Oliver10}), and possibly correspond to the active starburst progenitor phase of massive spheroidal galaxies (e.g.~\\citealt{Lilly99}).  However, several critical questions about SMGs remain unanswered.  The most fundamental of these is: \\textit{How is the immense luminosity of an SMG powered? -- by an intense single compact nuclear starburst (as seen in local ULIRGs) or by a more extended starburst (as seen in local normal star-forming galaxies)?}  One route to tackle this question is to ask what are the local analogues of SMGs and can these be used to understand the physics of the star formation activity within the distant SMG population?  Are SMGs simply `scaled up' ULIRGs, where the star formation occurs in a compact and highly obscured nuclear region -- potentially with a significant contribution from an active galactic nucleus (AGN; e.g.~\\citealt{Alexander05b})? Or does some fraction of the population resemble very scaled up `normal' star-forming infrared (IR) galaxies, which have far-IR emission arising from an extended, cooler component (termed `cirrus'; \\citealt{ERR03})? If we could empirically match subsets of the SMG population to local galaxies, then this would provide significant insights into the physical processes occuring within this important population. The synergy of \\textit{Herschel} and the Atacama Large Millimetre Array (ALMA) provides a novel opportunity to move beyond the first generation of detection experiments by beginning to examine the interplay between the interstellar medium (ISM) and the power source of high-redshift luminous dusty star-forming galaxies. \\smallskip  \\subsection{Far-infrared spectroscopy: a probe of SMG astrophysics}  The primary diagnostic lines of the ISM in the far-IR have traditionally included the fine-structure lines of [OIII]52, [NIII]57, [OI]63, [OIII]88, [NII]122, [OI]146, [CII]158\\,$\\mu$m, [CI]370\\,$\\mu$m, and [CI]610\\,$\\mu$m.  The relative strengths of these lines can be modeled as a function of the neutral gas cloud density and the cloud-illuminating far-UV radiation field intensity in photodissociation regions (PDRs; e.g.~\\citealt{Tielens85}; \\citealt{Kaufman99}; \\citealt{Wolfire03}; \\citealt{Meijerink05}), providing diagnostics of the physical conditions of the ISM. While much has been learned about the ISM conditions in Galactic star-forming regions and local galaxies through facilities such as the {\\it Infrared Space Observatory (ISO}; e.g.~\\citealt{Malhotra01}; \\citealt{Luhman98,Luhman03}; \\citealt{Negishi01}) and the Kuiper Airborne Observatory (KAO; e.g.~\\citealt{Crawford85}; \\citealt{Stacey91}; \\citealt{Poglitsch95}), due to previously poor instrument sensitivity and inaccessibility from the ground, dedicated surveys probing the ISM conditions in samples of high redshift dusty galaxies have only recently begun in earnest (e.g.~\\citealt{Stacey10}; \\citealt{Decarli12}) with the commissioning of purpose-built instrumentation like the redshift ($z$) and Early Universe Spectrometer (ZEUS; \\citealt{Stacey07}; \\citealt{HD09}).  For galactic and extragalactic systems, the `classical' PDR tracers are the brightest fine-structure emission lines [CII]158\\,$\\mu$m and [OI]63\\,$\\mu$m, acting as the primary coolants for the dense ($n\\gtrsim10$--10$^{5}$\\,cm$^{-3}$), warm (T$\\sim100$--1000\\,K), and neutral material.  Since [CII]158\\,$\\mu$m and [OI]63$\\mu$m correlate well with SFR and are sufficiently luminous (accounting for as much as $\\sim0.1$--1\\% of the bolometric luminosities in nearby galaxies and are thus much more luminous than molecular lines; \\citealt{Stacey91}), they have the potential to be used as powerful tracers of ISM conditions out to high-$z$.  However, \\textit{ISO} observations of local ULIRGs revealed that the [CII]158\\,$\\mu$m emission is about an order of magnitude lower relative to the far-IR (FIR) continuum flux than is seen in normal and starburst galaxies (e.g.~\\citealt{Malhotra01}; \\citealt{Luhman03}).  And this has recently been shown by \\citet{JGC11} using {\\it Herschel}-PACS to be a general aspect of far-IR fine structure lines (see Section~\\ref{discuss} and \\citealt{JGC11} for a discussion of a possible physical explanation). The potential weakness of [CII]158\\,$\\mu$m has implications for its use as a star formation tracer for high-redshift studies ($z>4$; e.g.~\\citealt{Maiolino05}; although cf.~\\citealt{Maiolino09}; \\citealt{DeBreuck11}). In order to make further progress in understanding the driver of star formation in luminous high-redshift SMGs in relation to galaxy populations in the local Universe, we need to begin compiling a similar set of (multi-transition) CO, [CII]158\\,$\\mu$m, and [OI]63\\,$\\mu$m data (which provide strong joint constraints on the PDR and star formation conditions such as the Hydrogen volume density, $n$, and the far-UV radiation field strength, $G_{0}$ in Habing units) for typical luminous star-forming SMGs at the epoch where their population peaks ($z\\sim$1--3).  \\subsection{This paper: New {\\it Herschel} observations of \\hbox{[OI]} 63\\,$\\mu$m in SMGs}  In the present paper we will be discussing observations performed with the ESA {\\it Herschel Space Observatory} \\citep{Pilbratt10}, in particular employing {\\it Herschel's} large telescope and powerful science payload to perform photometry and spectroscopy using the Spectral and Photometric Imaging REceiver (SPIRE; \\citealt{Griffin10}) and the Photodetector Array Camera and Spectrometer (PACS; \\citealt{Poglitsch10}). This paper presents the first attempt at compiling a census of the strength of [OI]63\\,$\\mu$m in far-IR luminous SMGs at the peak of star formation activity in the Universe ($z\\sim$1--2), enabling us to determine its suitability as a star formation tracer at high redshift.  These \\textit{Herschel} observations represent a key stepping stone to future ALMA studies of [CII]158\\,$\\mu$m and (multi-transition) CO to investigate the ISM physics of SMGs in detail, since [OI]63\\,$\\mu$m emission can only be measured in $1\\lesssim z\\lesssim 2$ SMGs with {\\it Herschel}-PACS, as the emission line falls outside of the ALMA bands for all but the highest redshifts ($z\\geq4$).  This paper is organized as follows: the sample selection, {\\it Herschel}-PACS spectroscopic observations and data reduction are described in Section~\\ref{obsdr}. In Section~\\ref{analysis}, we present the main analysis and results of the PACS spectrocopy and combine these measurements with L$_{\\rm FIR}$ estimates obtained from the combination of Large Apex Bolometer Camera (LABOCA; \\citealt{Siringo09}) and new {\\it Herschel}-SPIRE and PACS photometry.  We discuss the implications of our results in Section~\\ref{discuss}. Finally, our conclusions are given in Section~\\ref{conclusions}.  Throughout the paper we assume cosmological parameters of $\\Omega_\\Lambda=0.73$, $\\Omega_\\mathrm{m}=0.27$, and $H_\\mathrm{0}=71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ \\citep{Spergel03}. ",
        "conclusions": "We used {\\it Herschel}-PACS to target the [OI]63\\,$\\mu$m emission line in a sample of {\\it unlensed} high-redshift 870\\,$\\mu$m-selected $1.1<z<1.6$ SMGs for the first time.  Our sample originally consisted of nine statistically robust radio and/or mid-IR counterparts to the SMGs (providing precise positions for the SMGs) with high-quality spectroscopic redshifts.  We used high-resolution interferometric ALMA 870\\,$\\mu$m continuum data to verify the reliability of our original target selection.  The ALMA data have verified unambiguously that six of the SMG counterparts are bona fide submm-emitting sources; three of the SMG counterparts that were targeted here do not appear to be a source of submm emission and we have removed these from our sample.  We detect the [OI]63\\,$\\mu$m emission line in two of the SMGs, tentatively detect [OI]63\\,$\\mu$m emission in one SMG, and measure sensitive 3\\,$\\sigma$ upper limits to the line flux in the remaining three cases, secure in the knowledge that we have targeted the source of the submm emission and the correct redshift.  The [OI]63\\,$\\mu$m detections confirm the presence of a warm dense neutral ISM in these systems and indicate that [OI]63\\,$\\mu$m is an important PDR cooling line for SMGs. We find that the [OI]63\\,$\\mu$m--to--FIR ratio in SMGs (ranging from $\\simeq$0.5--1.5\\,per cent) does not follow the well-documented PDR cooling line `deficit' that has been observed for local ULIRGs and some local AGNs which are in a similar far-IR luminosity class to SMGs. This result suggests that local ULIRGs and AGNs are not good analogues of typical SMGs, and that SMGs appear to be more like `normal' starburst galaxies in terms of their star formation properties.  This offset in the [OI]63\\,$\\mu$m--to--FIR ratio compared to local (U)LIRGs has been noted previously for a small sample of SMGs in another classical PDR emission line, [CII]158\\,$\\mu$m, and together these results suggest that this could be a global property for all the bright PDR cooling lines for SMGs. We have taken advantage of a unique window of opportunity provided by {\\it Herschel} to verify that [OI]63\\,$\\mu$m is thus a reliable tracer of star formation in high-redshift star formation dominated galaxies such as SMGs.  These \\textit{Herschel} observations are a necessary stepping stone to future ALMA studies of [CII]158\\,$\\mu$m and CO emission in order to investigate the ISM physics of SMGs in detail."
    },
    "1208/1208.5698_arXiv.txt": {
        "abstract": "We present ultraviolet (UV) integrated colors of 44 Galactic globular clusters (GGCs) observed with the Galaxy Evolution Explorer (GALEX) in both $FUV$ and $NUV$ bands.  This data-base is the largest homogeneous catalog of UV colors ever published for stellar systems in our Galaxy. The proximity of GGCs makes it possible to resolve many individual stars even with the somewhat low spatial resolution of GALEX.  This allows us to determine how the integrated UV colors are driven by hot stellar populations, primarily horizontal branch stars and their progeny.  The UV colors are found to be correlated with various parameters commonly used to define the horizontal branch morphology. We also investigate how the UV colors vary with parameters like metallicity, age, helium abundance and concentration. We find for the first time that GCs associated with the Sagittarius dwarf galaxy have $(FUV-V)$ colors systematically redder than GGCs with the same metallicity. Finally, we speculate about the presence of an interesting trend, suggesting that the UV color of GCs may be correlated with the mass of the host galaxy, in the sense that more massive galaxies possess bluer clusters. ",
        "introduction": "The main contributors to the UV emission from any stellar system are the hottest stars. Indeed blue horizontal branch (HB) stars are well known to be among the hottest stellar populations in globular clusters (GCs) and contribute substantially to the UV radiation observed from old stellar systems (Welch \\& Code, 1972). Later on, two other sub-classes of post-HB stars were found to be important contributors to far-UV ($FUV$, $\\lambda \\sim 1500$--1600\\,\\AA) radiation (e.g., Greggio and Renzini 1990; Dorman et al. 1995, hereafter DOR95; Han et al. 2007). The hottest HB stars (extreme HB, EHB) have such a small envelope mass that most of their post-He-core burning phase takes place at high effective temperature (\\teff), during the so called \"AGB-manqu\\'e phase\", and these stars never return to the asymptotic giant branch (AGB). Another group of UV-bright stars is that of post-early AGB stars, which after a brief return to the AGB, spend the bulk of their helium shell burning phase at high \\teff. In systems with only red HB a small floor level of $FUV$ is provided by post-AGB stars, which evolve to the AGB phase with an higher envelope mass where they undergo thermal pulses and eventually lose their envelopes moving at higher temperatures at constant luminosity. The relative contributions of the various types of stars and the factors that might lead to larger or smaller populations of UV-bright stars have remained an open question (Greggio and Renzini 1990; DOR95; Lee et al.  2002; Rich et al. 2005; Sohn et al. 2006).  In distant extragalactic systems one can ordinarily observe only the integrated light of unresolved stellar populations, from which the hope is to gain knowledge about the underlying stellar population. Galactic globular clusters (GGCs) play an important role in understanding the integrated UV colors of extragalactic systems, especially the so called \"UV-upturn\" observed in the spectral energy distributions of elliptical galaxies (Code \\& Welch 1979; de Boer 1982; Bertola et al. 1982; Greggio \\& Renzini 1990; O'Connell 1999). First of all, GCs are the closest example in nature to a single stellar population (SSP): a system of coeval stars with similar chemical composition\\footnote{Although there is now a general consensus that the formation of GGCs may have been more complex than previously thought (based for example, on the detection of chemical inhomogenites in light elements, Carretta et al. 2009a, and the existence of multiple populations, Piotto 2009), their stellar content has been found to be quite homogeneous in terms of iron abundance. Indeed only two GC-like stellar systems have been found to host multiple populations with significant ($>0.5$ dex) spread in the iron abundance and (possibly) age: $\\omega$ Centauri (Norris et al. 1996, Lee et al. 1999, Ferraro et al. 2004) and Terzan~5 (Ferraro et al. 2009).}.  Moreover GGCs span a large range of metallicities, a small range of ages, and perhaps some range of helium abundance. Hence they can be used to test the stellar evolution theory, which in turn is one of the basic ingredients of the models used to interpret the integrated light of distant galaxies. GGCs are relatively nearby objects (more than $\\sim90\\%$ are located at distances $r<30$\\,kpc), so their populations can be easily resolved.  With typically more than 100,000 stars, even relatively short-lived evolutionary stages are sampled.  We can directly observe the properties of individual stars and measure the population ratios for objects in different evolutionary stages. In particular we can study the impact of hot and bright populations (as the AGB-manqu\\'e stars) on the integrated UV light of GGCs and then use them as crucial local templates for comparison with integrated properties of distant extragalactic systems. In fact comparing features in the color-magnitude diagrams (CMDs) of well known and resolved GGCs with integrated quantities can lend important {\\it model independent} insights into the nature of extragalactic systems.  Integrated UV photometry of GGCs has previously been obtained by the \\emph{Orbiting Astronomical Observatory} (OAO~2; Welch \\& Code 1980), \\emph{Astronomical Netherlands Satellite} (ANS; van Albada, de Boer \\& Dickens 1981), \\emph{Ultraviolet Imaging Telescope} (UIT; Hill et al. 1992; Landsman et al. 1992; Parise et al. 1994; Whitney et al. 1994), and \\emph{International Ultraviolet Explorer} (IUE; Castellani \\& Cassatella 1987). Using a large, but heterogeneous, collection of data obtained with some of these telescopes along with population synthesis models, DOR95 showed how the UV colors varied with parameters like metallicity and how they compare with elliptical galaxies. They showed how the UV colors of GGCs could plausibly be produced by hot HB stars and their progeny.  At the time of DOR95, UV photometry of individual stars in GGCs was available for only a few clusters. That situation has changed dramatically.  Our group alone has already published HST UV photometry for a dozen of GGCs (see Ferraro et al.  1997, Ferraro et al. 1998, Ferraro et al.  1999, Ferraro et al.  2001, Ferraro et al. 2003, Lanzoni et al. 2007, Dalessandro et al.  2008, Rood et al. 2008) and obtained data for an additional 32 clusters in HST Cycle 16S (GO11975, PI: Ferraro; see Contreras et al. 2012, Sanna et al. 2012 and references therein).  More recently, we have secured observations of 44 GGCs during three observing cycles with the \\emph{Galaxy Evolution Explorer} (GALEX). This is the largest homogeneous sample ever collected for GGCs in UV so far.  In Schiavon et al. (2012; hereafter Paper I) we presented photometry and CMDs for these clusters.  Here we present integrated UV magnitudes and colors (\\S~2) for each cluster, and we describe (\\S~3) how they are affected by HB class, metallicity, age, and possibly structural parameters (mass, density, central relaxation time, etc.). In \\S~4 we compare our data with observations of GCs in M31 and M87. ",
        "conclusions": "As part of a project aimed at studying the properties of hot stellar populations in the Milky Way GCs (see Paper I), we have presented UV integrated colors obtained with GALEX for 44 clusters spanning a wide range of metallicities ($-2.5<\\feh<-0.4$), HB morphologies, structural and dynamical parameters. This represents the largest homogeneous catalog of UV photometry ever built for GGCs.  We compared the behavior of UV colors with several parameters characterizing the morphology of the HB. As expected, there are general correlations, in particular between $(FUV-V)_0$ and the $(B2-R)/(B+V+R)$ parameter defined by Buonanno et al. (1993; 1997) and the HB temperature extension \\teff\\ defined by Rec\\'\\i o Blanco et al. (2005). There is also a significant correlation with the $\\Delta (V-I)$ parameter (Dotter et al. 2010), but, as expected, this parameter is insensitive to HBs with extreme blue extensions.  In all color combinations, the bluest clusters are those in the intermediate metallicity regime ($-1.5<\\feh<-1$). This is in agreement with DOR95 and with the L$_{\\rm t}$ parameters measured by Fusi Pecci et al. (1993).  In the $(NUV-V)_0$-$\\feh$ plane, clusters more metal rich than $\\feh\\sim-1.5$ show a clear and significant linear correlation with cluster becoming redder as metallicity increases, while there is an opposite trend in the metal-poor regime. The combinations of colors involving the $FUV$ appear more scattered, but a reasonable and similar correlation with metallicity at ($2.5\\sigma$ level) has been also found in these cases. All the clusters suspected to be connected with the Sagittarius dwarf spheroidal (NGC~4590, NGC~5053, NGC~5466, Arp~2 and Terzan~8) are typically $\\sim 1.5$\\,mag redder in $FUV$ colors than systems with similar iron content. No appreciable differences are found in $(NUV-V)_0$. Studies from different groups suggest that Sagittarius clusters and their galactic counterparts are coeval, while GGCs are on average more He-rich than the Sagittarius sub-set. This would tentatively be interpreted as due to a different environment in which they formed.  With the aim of showing how sensitive ages derived from UV colors may be to assumptions about helium abundance, we compared our colors with evolutionary models of SSP by Lee et al. (2002).  The color-metallicity distribution of GGCs can be reproduced by assuming an average age of $\\sim 12$ Gyr, with a spread of about $\\pm 2\\,\\gyr$. Alternatively, in the framework in which some GCs have experienced self-enrichment from material ejected from AGBs or fast-rotating massive stars (Ventura \\& D'Antona 2008; Decressin et al. 2007), the color spread is consistent with different He content. In particular, we show that an overabundance of helium ($\\Delta Y(R') \\sim 0.05 $ ) can mimic an age difference $\\Delta t \\sim 2\\,\\gyr$.  The UV colors of GGCs are consistent with those obtained by GALEX for M31 clusters (Rey et al. 2007; Kang et al. 2011), at least in the intermediate/low metallicity regime. At $\\feh>-1$, M31 GCs are systematically bluer by 1--2 mag, behaving like massive MW clusters with bimodal HBs, such as NGC~6388 and NGC~6441. As already noticed by Sohn et al. (2006), M87 GCs are on average bluer than GGCs.  This might be the signature of different chemical abundances impressed on the \"integrated\" properties of GCs. In particular we speculate that He abundance may be correlated with the mass of the host galaxy, being higher in GCs belonging to higher mass galaxies.\\\\      {\\it The authors dedicate this paper to the memory of co-author Bob Rood, a pioneer in the theory of the evolution of low mass stars, and a friend, who sadly passed away on 2 November 2011.}\\\\ We thank the anonymous referee for the useful comments and suggestions. The authors warmly thank M. Bellazzini and F. Fusi Pecci for useful discussions that improved the presentation of the results. We are grateful to M. Catelan for providing us informations about HB parameters. This work is based on observations made with the NASA Galaxy Evolution Explorer. GALEX is operated for NASA by the California Institute of Technology under NASA contract NAS5- 98034.  This research is part of the project COSMIC-LAB funded by the European Research Council (under contract ERC-2010-AdG-267675).  E.D. thanks the hospitality and support from Gemini Observatory where most of this work was developed, during two extended visits in 2009 and 2010, and a shorter visit in 2011. R.P.S. acknowledges funding by GALEX grants $\\#$ NNG05GE50G and NNX08AW42G and support from Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., on behalf of the international Gemini partnership of Argentina, Australia, Brazil, Canada, Chile, the United Kingdom, and the United States of America.  S.T.S. acknowledges support for this work from the GALEX Guest Investigator Program under NASA grant NNX07AP07G.      \\newpage"
    },
    "1208/1208.0242_arXiv.txt": {
        "abstract": "{In the exoplanetary era, the Kepler spacecraft is causing a revolution by discovering thousands of new planet candidates. However, a follow up program is needed in order to reject false candidates and to fully characterize the bona-fide exoplanets.} {  Our main aims are: 1./ Detect and analyze close companions inside the typical Kepler PSF to study if they are the responsible of the dim in the Kepler light curves, 2./ Study the change in the stellar and planetary parameters due to the presence of an unresolved object, 3./ Help to validate those Kepler Objects of Interest that do not present any object inside the Kepler PSF and 4./ Study the multiplicity rate in planet host candidates. Such a large sample of observed planet host candidates allows us to do statistics about the presence of close (visual or bounded) companions to the harboring star. } {We present here Lucky Imaging  observations for a total amount of 98 Kepler Objects of Interest. This technique is based on the acquisition of thousands of very short exposure time images. Then, a selection and combination of a small amount of the best quality frames provides a high resolution image with objects having a 0.1 arcsec PSF. We applied this technique to carry out observations in the Sloan i and Sloan z filters of our Kepler candidates.} {We find blended objects inside the Kepler PSF for a significant percentage of KOIs. On one hand, only 58.2 \\% of the hosts do not present any object within 6 arcsec. On the other hand, we have found 19 companions closer than 3 arcsec in 17 KOIs. According to their magnitudes and $i-z$ color, 8 of them could be physically bounded to the host star.} {} ",
        "introduction": "It is not so long when the discovery of extrasolar planets was just utopian. However, after the first discovery of an exoplanet orbiting a main sequence star \\citep{mayor95,marcy96}, the scientific community has discovered and confirmed hundreds of these objects around other stars. In this context, the Kepler Space Telescope  has detected a new large sample of planet host candidates by continuously observing more than 150000 stars in a specific field of view (RA=19h 22m 40s DEC=+44 30' 00''). During the first five quarters of observations (i.e, $\\approx 4.5$ months) the Kepler Team collected on its second public release a total amount of 997 planet host star candidates \\citep[second public release,][]{borucki11}.  However, to date, less than 5\\% of these candidates have been confirmed. After the candidates selection, there is a mandatory step to reject false positives before attempting any high accurate (difficult and expensive) radial velocity measurements. Since the Kepler effective point-spread function is very large (6-10 arcsec, depending on the particular target) and its pixel size is about 4 arcsec, several background objects could be blended by the host candidate (called Kepler Object of Interest, hereafter KOI). Bounded or not, these objects clearly affects the star and planet parameters such as the planet-to-star radius ($R_p/R_*$), semi-major axis ($a/R_*$), impact parameter ($b$) or even the planetary mass ($M_p$).  The presence of a secondary star could lead to the definite rejection of the candidate, as an example, \\citep[see][]{odonovan06} In fact, there are several configurations that can mimic an exoplanet transit around its host star. The most relevant are: A./ a small substellar object transiting the other component of the binary system (since the smallest stars and brown dwarfs have the same size as Jupiter), B./ a stellar binary blended with a background star, C./ a grazing binary, which has not been ruled out by additional photometry or spectroscopy, D./ a background eclipsing binary blended by the light of the KOI, E./ a larger planet transiting a background star blended by the main target (this actually would not strictly be a false positive since there is a planet in the system but it would be in the sense of rejecting the brightest central star as a planet host) and  F./ a long-term spot. There would be a 'G-case' where the main target actually hosts a planet but with a blended background companion or a non transiting bounded companion. These configurations would lead to a change in the planet properties, as it was said before.  Some of these configurations might be ruled out by the automatic pipeline implemented by the Kepler Team \\citep{jenkins10}. While cases A and B are well rejected by this pipeline and an individual study of the light curves done by the team itself \\citep{borucki11},  low-resolution spectroscopy clearly reject the C configuration. However,  D, E and G cases are the main sources of false positives in the sample of transiting planet candidates. More specifically, case G clearly shows the need for an intense high resolution imaging follow-up program to validate the planetary nature of the transients. Due to Kepler long base-line, we expect few or no case F.  Theoretical studies in regard to the false positive probability of Kepler candidates conclude that obtaining high resolution images down to 1-2 arcsec is crucial to confirm the planets and its physical properties. As an example, an earth-size planet transiting a faint star might have a false positive probability greater that 20\\% if it lacks high-resolution imaging,  which could potentially be decreased to less than 2\\% with a high-resolution image \\citep{timothy11}. Several authors have acquired this kind of observations for other planet host candidates finding important changes in the planet-star properties. For instance, \\cite{daemgen09} found stellar companions to 3 stars harboring planets. As a consequence, the updated values of the physical parameters changed about 2\\% with respect to the previous ones.    However, even if the planet is confirmed, its formation and evolution scenarios (including the migration) require accurate description of the effect of bounded stellar companions. The vast majority of planets found in multiple systems are actually S-type, meaning that the planet is orbiting one of the components of the multiple system playing the secondary star the role of a gravitational perturber \\citep{kley10}. The presence of these secondary objects difficult the planet formation since they interact dynamically with the elements of the system and produce an extra heating of the protoplanetary disk. All these factors may imply important changes in the planetary architecture and exoplanet properties with respect to those formed around single stars. For instance, \\cite{eggenberg04} found a statistical segregation in the planet mass for those planets with orbital periods shorter than  40 days around single and multiple systems.   In this paper, we present the first statistical study of multiplicity on the Kepler candidates. A set of high resolution images obtained with the Lucky-Imaging technique \\citep{law06} in the 2.2 m telescope in Calar Alto Observatory (Almer\\'ia, Spain) with the AstraLux instrument were acquired.  This technique allows us to obtain diffraction limited observations with the best seeing conditions in the SDSSi band (see the F\\'elix Hormuth PhD dissertation). A total amount of 98 KOIs (i.e., about a 10\\% of the candidates listed in the Kepler second public release of the Kepler Team) have been pointed and studied.   In section \\S2 we will explain the observations, image processing and data extraction from the raw images, regarding: sample selection (\\S2.1), data acquisition (\\S2.2), data reduction and photometric calibration (\\S2.3), including the astrometric corrections, and a study of the sensitivity curves and detectability of our observations. Results based on these observations will be presented in section \\S3. We also perform spectral typing of the detected companions (\\S4.1) and give some clues about the possible gravitational bound between some of them (\\S4.2). In section \\S4.3 we will discuss the implications on calculated planetary parameters from the Kepler light curves regarding the presence of a blended star. Some particularly interesting cases will be studied in \\S4.4 and final conclusions of this work will be presented in section \\S5. ",
        "conclusions": "We have started a high resolution imaging follow-up for the Kepler sample of planet host candidates. The main goal of this survey is to provide additional constrains for the confirmation of the planetary nature  of these candidates and identify those that are possible false positives. A total amount of 98 KOIs (out of the 997) from the second release of the Kepler Team have been properly observed by using the Lucky Imaging technique with the AstraLux instrument at the 2.2m telescope at Calar Alto Observatory.  Our main results show that  the 58.2\\% of the KOIs are actually isolated in terms of not having any visual or bounded companion at less than 6 arcsec. In other words, the 41.9\\% of the candidates present close objects. This is an important result in terms of: (a) False positive rate determination, since it points directly which objects have stellar companions that can be mimicking a planet transit detected by Kepler, being then the highest priority for a deeper follow-up with ground-based telescopes to determine the nature of this transit; (b) Updating the planet properties, since as we have shown in the \\textit{Discussion} section, they depend on the brightness of the host star. We warn that orbital and physical parameters of the 23 planets orbiting the 17 KOIs with close stellar companions should be revised. (c) Estimating the binary rate in planet host stars. According to their position on a $M_i$ vs. $i-z$ color-magnitude diagram, we have shown that between 6 and 9 of the close companions could be actually bounded to the host star due to their position over our empirical ZAMS. Their distances agree with an S-type binary although still more observations are necessary to confirm both the planet and the binary in all cases. KOIs 0379B, 0658B, 0641B, 0645B and 0298B clearly lie over the ZAMS which suggests a simultaneous formation together with the primary star. Moreover, KOIs 0433B, 0401B, 0592B and 0703B could also lie inside the error bars of the ZAMS but more work should be done to confirm this result. If confirmed, it would imply a lower limit on the observed binary frequency of 6.2-9.2\\% Regarding the medium-distance companions (3-6 arcsec), we concluded that only one of them (KOI-0623B) is possibly bounded although we should flag this result due to the large errors in the distance estimates.  Finally, we have provided accurate astrometric positions and $i$ magnitudes for the close and medium distance companions which could be used to re-compute planet-star parameters in those KOIs affected by the light of the companion. These results add more constrains for theoretical works regarding false positive probabilities for the particular objects studied in this paper."
    },
    "1208/1208.2071_arXiv.txt": {
        "abstract": "The joint use of high-resolution data from PROBA2 and SDO satellites and LASCO/SOHO coronographs enabled us to examine early stages of initiation and propagation of six limb CMEs registered in June 2010 - June 2011. For five events under consideration, the CME initiation is marked by filament (prominence) eruption or by a loop-like structure having another nature.  Subsequently, several loop-like structures having higher brightness and following each other at different velocities appear in the region of the CME initiation. The CME frontal structure is formed by these loop-like structures. The time-dependent velocities and acceleration of the ejection front have been obtained for all CMEs under consideration. We have drawn a conclusion about the possible existence of two CME types dependent on the time profile of their velocity. The first CME type comprises the ejections whose velocity decreases abruptly by more than 100 km/s after having reached the maximum; it thereupon passes to slow deceleration. The second CME type is formed by the ejections whose velocity varies insignificantly after reaching the maximum. The CME angular size is shown to increase up to threefold at the initial stage of propagation; it increases twofold 3.5-11 minutes after the first measurement of this parameter. When considering 3 CMEs, we see that their broadening exceeds their extension in the longitudinal direction during a certain period of time at the initial propagation stage. ",
        "introduction": "\\label{S-Introduction}   Physical mechanisms of generation of coronal mass ejections (CMEs) still remain unclear in many respects \\cite{Forbes}. This burning issue can be addressed due to the use of data with high temporal and spatial resolution that will allow us to study initiation and initial propagation stage of CMEs. In 2010, Solar Dynamics Observatory (SDO) \\cite{Pesnell} was launched with the on-board Atmospheric Imaging Assembly (AIA) instrument \\cite{Lemen} having unique characteristics. AIA enables us to observe the Sun in the wavelength range from extreme ultraviolet ($\\lambda=9.4 \\; nm$ ) to ultraviolet ($\\lambda=170.0 \\; nm$ ) with high temporal (10 seconds) and spatial resolutions (0.6 arcsec pixels). Consequently, it will be possible to reveal regularities of the CME generation and the initial phase of the CME propagation.  The additional use of data from the PROBA2 (Project for On-Board Autonomy) satellite launched in 2009 will contribute to this study, too (\\url{http://proba2.oma.be/index.html/about/}). The on-board instrument SWAP \\\\ \\cite{Katsiyanni} observes the Sun in extreme ultraviolet at 17.4 \\textit{nm}, whereas LYRA (\\url{http://proba2.oma.be/index.html/science/}) monitors the solar irradiance in the wavelength range from soft X-rays to ultraviolet. Though the temporal (1 minute) and spatial (3.17{''}) resolutions of SWAP are worse than those of AIA instruments, it has a big advantage: its field of view is significantly larger (54 arcmin as against 41 arcmin (SDO)).  Establishment of experimental regularities of the CME propagation immediately after its initiation is of importance when studying CME features. Over the last decade, several studies were devoted to different aspects of the initial stage of the CME propagation, using data from TRACE, SOHO, STEREO, GOES satellites (see papers \\cite{Gallagher, Zhang, Maricic09}). First of all, kinematics of the CME propagation was examined. Three periods of the initial propagation of CMEs were established \\cite{Zhang}: slow increase in velocity within several tens of minutes, rapid increase in velocity (the main CME acceleration) within the period from several minutes to several tens of minutes (sometimes up to several hours) when the acceleration amplitude reaches the value of several $km/s^2$ in some ejections. The third period is the propagation phase of CMEs (i.e., the 'quiet' propagation of CMEs with an acceleration that is small or moderate in absolute magnitude).  A close connection of the CME acceleration with energy release was established in the CME-related solar flares \\cite{Zhang, Maricic07, Temmer08, Temmer10}. There was a positive correlation between the duration of the main CME acceleration tACC and the time of increase tSXR in intensity of soft X-rays ISXR from the CME-related flare area up to maximum values \\cite{Zhang, Maricic07}. The CME acceleration $a(t)$ and hard X-ray emission $I_{HX}(t)$ (or its related derivative $dI_{SXR}/dt$) from the flare area were shown to be often synchronised \\cite{Maricic07, Temmer08, Temmer10}. An inverse correlation dependence of acceleration $V_{max}/t_{acc}$ on $t_{acc}$ (or $V_{max}/t_{SXR}$ and $t_{SXR}$) was established \\cite{Zhang, Vrsnak}.  Due to unique capabilities of the instruments on board SDO, new results on the CME initiation and its initial propagation stage have been obtained \\cite{Patsourakos2}. With a CME as an example, we have shown that the CME initiation has three stages: the stage of slow self-similar expansion, short-term strong hyperexpansion in the transverse direction, and one more stage of self-similar expansion in the SDO field of view. There is also a strong short-term CME acceleration with a subsequent deceleration. Note that the conclusion about the existence of a period of the CME hyperexpansion in the transverse direction was drawn from PROBA2 data later in \\cite{Fain}.  In this paper, we go on studying the CME initiation during the new solar cycle, using SDO/AIA data. We also examine the initial stage of the CME propagation with the use of data from the SDO, PROBA2 and SOHO spacecraft. The goal of this study is to get new experimental results concerning the CME kinematics and time variations of their geometric parameters that could be used to construct adequate models of the CME generation and propagation at the initial stage. ",
        "conclusions": "% \\label{Discussion and conclusions} Modern models of the CME formation are provided by a limited amount of experimental data on CMEs at early stages of their initiation and propagation \\cite{Forbes}. The aim of this paper is to establish new regularities of the CME formation and of the initial stage of its propagation in order to use them for creating an adequate model of the CME initiation. This is possible due to data from instruments AIA aboard SDO that have unique characteristics, as well as due to joint data from PROBA2/SWAP, SOHO/LASCO and GOES.  We have established physical and morphological features of formation and initial stage of propagation of six CMEs: \\begin{enumerate} \\item event starts from the eruption of filament or loop-like structure having another nature.  This has been reliably established for five events of 13 June 2010, 11 February 2011, 27 March 2011, 27 April 2011, and 07 June 2011; \\item soon after the filament eruption several overlying loops occur and start moving at different velocities, thus forming the CME frontal structure (for all the events under consideration); \\item during the four events (13 June 2010, 11 February 2011, 27 March 2011, 07 June 2011), velocity of the CME front decreases abruptly by more than 100 km/s after having reached the maximum; it then varies insignificantly in the LASCO field of view. We suppose that such events form a special CME type. Velocity of two CMEs (08 March 2011 and 27 April 2011) decreases slightly after reaching the maximum value. Probably these CMEs represent another type of CMEs. According to the conclusion in \\cite{Zhang}, this pattern of velocity variation is typical for CMEs. Findings in \\cite{Temmer08, Temmer10, Patsourakos2, Fain, FainZ} and in this paper prove that the most typical is the profile of the CME velocity where velocity decreases abruptly by more than 100 km/s after having reached the maximum. \\end{enumerate}  The formed CME frontal structure is not the same in different spectral lines. The line $\\lambda=17.1\\;nm$ is formed at a lower temperature than the line $\\lambda=21.1\\;nm$. The shift of boundaries of the CME frontal structure along the radius implies that temperature of outer regions of the CME frontal structure is higher than that of inner regions.  Variation in the CME angular size $ 2\\alpha (t)$ demonstrates distribution and time variations in forces that influence the ejection boundary both outside and inside CME. The conclusion concerning the increase in the CME angular size $2\\alpha$ between CME locations in the lower corona and in the LASCO field of view has been drawn in cite{Plunkett}. However, no data on typical temporal and spatial scales of such variations have been obtained. We managed to observe time variations in angular sizes of four CMEs in the SDO field of view. Value $2\\alpha$ was found to increase up to threefold at the initial stage of propagation. Ratio of the CME angular size in the LASCO field of view to the first measured value $2\\alpha$ could be $\\approx$ 5 (for CME of 13 June 2010) and more. The typical temporal scale of doubling of the CME angular size after the first measurement of $2 \\alpha$ was $\\approx$ 3.5-11 minutes. In the frame of a simple CME model at the initial stage ('bubble', see \\cite{Patsourakos2}), increase in the CME angular size with time evidences excess of internal force's influence on the CME boundary over external force's influence.  Study of ratio of the limb CME's longitudinal size to its transverse size $d_H/d_W$ provides solutions to several problems. First of all, this ratio is equivalent to the ratio of the CME velocity along its axis to the CME expansion velocity in the transverse direction. Knowledge of formation laws of this ratio allows us to estimate, for instance, propagation velocity of halo CMEs along their axes in three-dimensional space, by measuring the CME expansion velocity in the transverse direction (see discussion on this issue, based on LASCO data, in \\cite{Michalek}). Besides, this ratio can be used to determine whether or not the CME expansion is self-similar. Comparison between CME propagation parameters and theory of self-similar expansion of CMEs \\cite{Uralov} provides a more accurate estimate of self-similar CME propagation.  If ratio $d_H/d_W$ remains constant over time, this implies conservation of the CME form (self-similarity) over time.  Note that the parameter we have chosen to determine the CME expansion differs from that used in \\cite{Patsourakos2} to perform the same analysis.  That parameter was referred to as 'aspect ratio' (ratio of the height of centre of circle $H_C$, which outlined well most of the CME boundary, to the radius of this circle $R_F$). Height $H_C$ was measured relative to the solar limb. We used the parameter closest to that in \\cite{Michalek}. It can be shown that two these parameters should be connected by ratio $H_C/R_F \\approx 2d_H/d_W-1$ for the event of 13 June 2010. When deriving this formula, we ignored deviation of location of the CME-related flare from the limb and from the centre of the CME basis.  Fig. 8 (c) compares dependence of $H_C/R_F$ on time \\cite{Patsourakos2} with our dependence $[d_H/d_W](t)$ for the event of 13 June 2010. We see that $H_C/R_F$ variation over time is qualitatively close to dependence $[d_H/d_W](t)$. In the first approximation, relation between the two parameters is described by their above-stated ratio. Noteworthy is the fact that, starting from instant $t=$ 05:39:00, shape of the CME boundary does not correspond precisely to circle (e.g., compare the visible CME boundary at $t = $05:40:11 in our figure with circle in Fig. 3 from \\cite{Patsourakos2}).  In conclusion, let us note that the use of data with high temporal and spatial resolution made it possible to obtain information about initiation and initial propagation stage of CMEs. However, a great deal needs to be done before initiation and development of CMEs become completely understood. In our next paper, we will try to relate the differences in $V(t)$ profiles established for 2 CME types to the peculiarities of the magnetic field and other features of active regions wherein CMEs initiate.    \\begin{acks} We are grateful to SDO/AIA, PROBA2/SWAP, SOHO/LASCO, \\\\ GOES and RHESSI teams for making their data available, and to V.V. Grechnev for fruitful discussions and his help in processing SDO data. This work was partially supported by grant No. 02.740.11.0576 under the Federal Programme 'Scientific and Scientific-Pedagogical Staff of Innovative Russia'. \\end{acks}   \\mbox{}~\\\\"
    },
    "1208/1208.0597_arXiv.txt": {
        "abstract": "\\noindent This is the first in a series of papers in which we measure accurate weak-lensing masses for \\clusterfields\\, of the most X-ray luminous galaxy clusters known at redshifts $0.15\\lesssim z_{\\rm Cl}\\lesssim 0.7$, in order to calibrate X-ray and other mass proxies for cosmological cluster experiments. The primary aim is to improve the absolute mass calibration of cluster observables, currently the dominant systematic uncertainty for cluster count experiments. Key elements of this work are the rigorous quantification of systematic uncertainties, high quality data reduction and photometric calibration, and the ``blind'' nature of the analysis to avoid confirmation bias. Our target clusters are drawn from X-ray catalogs based on the ROSAT All-Sky Survey, and provide a versatile calibration sample for many aspects of cluster cosmology. We have acquired wide-field, high-quality imaging using the Subaru and CFHT telescopes for all \\clusterfields\\, clusters, in at least three bands per cluster. For a subset of 27 clusters, we have data in at least five bands, allowing accurate photometric redshift estimates of lensed galaxies. In this paper, we describe the cluster sample and observations, and detail the processing of the SuprimeCam data to yield high-quality images suitable for robust weak-lensing shape measurements and precision photometry. For each cluster, we present wide-field three-color optical images and maps of the weak-lensing mass distribution, the optical light distribution, and the X-ray emission. These provide insights into the large-scale structure in which the clusters are embedded. We measure the offsets between X-ray flux centroids and the Brightest Cluster Galaxies in the clusters, finding these to be small in general, with a median of 20 kpc. For offsets $\\lesssim 100$ kpc, weak-lensing mass measurements centered on the Brightest Cluster Galaxies agree well with values determined relative to the X-ray centroids; miscentering is therefore not a significant source of systematic uncertainty for our weak-lensing mass measurements. In accompanying papers we discuss the key aspects of our photometric calibration and photometric redshift measurements (Kelly et al.), and measure cluster masses using two methods, including a novel Bayesian weak-lensing approach that makes full use of the photometric redshift probability distributions for individual background galaxies (Applegate et al.). In subsequent papers, we will incorporate these weak-lensing mass measurements into a self-consistent framework to simultaneously determine cluster scaling relations and cosmological parameters. ",
        "introduction": "\\label{sect:intro}  The formation of cosmic structure depends sensitively on the mass and energy content of the Universe, and the physical nature of dark matter and dark energy.  Galaxy clusters are the most massive gravitationally bound structures, sitting at the largest nodes of the cosmic web.  As such, their number density, baryon content, and evolution are sensitive probes of cosmological parameters, in particular the amplitude of matter fluctuations ($\\sigma_8$), the mean matter and dark energy densities ($\\Omega_{\\rm m}$ and $\\Omega_{\\rm DE}$), and the dark energy equation of state parameter ($w$) \\citep[for a recent review, see][]{aem11}.  The idea of ``counting clusters'' as a way to test cosmology has existed for decades \\citep[e.g.,][]{kai84,hea91}.  The discovery of massive clusters at high redshifts \\citep{dvg98,baf98} provided supporting evidence for a low matter density Universe, and presaged the discovery of dark energy from Type Ia supernovae studies \\citep{rfc98,pag99}.  Cluster counts paved the way in determining the now accepted value of $\\sigma_8\\sim0.8$ \\citep[e.g.][]{brt01,gbc03}. Recently, measurements of the evolution of the cluster number density have provided some of the most precise and robust constraints on dark energy \\citep{vkb09,mar10}, as well as departures from General Relativity on cosmological scales \\citep{ram10,rba12,svh09}, and the species-summed neutrino mass \\citep{mar10c,rvj10}.  A fundamental challenge for cluster count experiments is that the survey observations do not measure cluster masses directly, but rather a property that correlates with cluster mass, typically with significant associated scatter. For X-ray surveys, the standard survey observable is the X-ray flux, which with the cluster redshift gives the X-ray luminosity; for optical red-sequence finders, survey measures are typically based on optical richness; and for millimeter surveys the typical observable is the Sunyaev-Zel'dovich (SZ) flux. In order to reconstruct the underlying mass function, the scaling relation between the survey observable and cluster mass, as well as the scatter in this relation as a function of mass and redshift, must be measured.  This process can be improved if, for a representative subsample of the survey clusters, one can also obtain deeper follow-up measurements of precise mass proxies with lower systematic scatter \\citep{mar10,mae10,vbe09,vkb09}. The subsample re-observed need not be large in order to bring a substantial boost in constraining power \\citep{mar10,wrw10}.   \\subsection{The role of mass proxies}  X-ray observations provide a critical element of this work, offering several observables that are straightforward to measure and which correlate tightly with true cluster mass. For example, the temperature of the intracluster medium, $T_{\\rm X}$, traces cluster mass with a scatter of 10--15\\%, far better than the total X-ray luminosity (scatter $\\sim$40\\%). Other X-ray proxies such as gas mass, $M_{\\rm gas}$, thermal energy, $Y_{\\rm X} \\,(=M_{\\rm gas} T_{\\rm X})$, and center-excised X-ray luminosity provide comparable or possibly even lower scatter \\citep{ars08,kvn06,mau07,mae10}.  However, even for these low-scatter mass proxies, the absolute scaling with true cluster mass must also be determined, accurately and robustly. For X-ray data, under the assumptions of hydrostatic equilibrium and spherical symmetry, one can relate the observed gas density and temperature profiles to the underlying mass profile.  Yet even for the most dynamically relaxed clusters, and at optimal measurement radii ($r\\sim r_{2500}$), hydrostatic X-ray mass estimates are expected to be biased at the 5--10\\% level due to non-thermal pressure support from residual gas bulk motion and other processes \\citep{nvk07,rmm12}. For less relaxed systems, and for measurements at larger radii ($r\\gtrsim r_{500}$), the biases in hydrostatic measurements can be significantly worse \\citep[20--30\\%,][]{nvk07}. This uncertainty in the absolute mass scaling is currently the dominant systematic uncertainty in the constraints on $\\sigma_8$ from cluster counts \\citep{mar10,vkb09,rwr10,sta11,bhd11}. In order for future surveys to access their full constraining power, it is imperative to calibrate these mass proxies to within 5\\% and over the entire mass and redshift range of interest \\citep{wrw10}.  \\subsection{Weak-lensing mass measurements as calibrators for cluster masses}  The most promising method currently capable of absolutely calibrating mass measurements for statistical cluster samples is cluster weak gravitational lensing. Weak-lensing mass measurements do not require a baryonic tracer, but directly measure the total gravitating matter. For individual clusters, weak lensing is inherently noisy since the intrinsic ellipticity distribution of galaxies is broad and lensing measurements are sensitive to all structure along the line of sight. Utilizing cluster weak-lensing mass measurements for precision cosmology requires a thorough understanding of the systematic biases involved. Since the shear induced on a background galaxy depends on the cluster mass, the ratios of angular diameter distances between the observer, cluster and source, and cosmology, there are three possible sources of systematic uncertainties. Observationally, biases in the shear measurements and in the redshifts of background galaxies translate to biased mass measurements.  Even in the absence of observational biases, systematic uncertainties may arise from the assumptions made to relate the measured lensing signal to an intrinsic cluster mass.  Lensing inherently measures projected, 2D masses; however, to compare these to the halo mass function, they need to be related to 3D masses. The most common method to do so is to fit spherically symmetric density models \\citep[such as the NFW profile, ][]{NFW97} to the measured shear profiles. Adopting a profile shape has the added advantages that it breaks the mass-sheet degeneracy, and that significantly fewer galaxies are required compared to non-parametric mass reconstruction. (Note that the aperture mass method, which also assumes spherical symmetry but does not directly fit a specific profile, still requires a profile assumption at large radii to break the mass-sheet degeneracy.) However, because clusters are generally triaxial, the assumption of spherical symmetry leads to over-/underestimates of the mass if the cluster major axis is aligned along/perpendicular to the line of sight \\citep{cok07,mrm10}.  Mass in the infall region of clusters (e.g., filaments and infalling groups) and/or unassociated structures along the line of sight can similarly bias individual mass measurements \\citep{hoe01,hoe03}.  Quantifying the expected scatter due to these sources, as well as any expected bias due to the profile assumption, can be achieved straightforwardly from cosmological N-body simulations, by applying the same mass measurement methods to the simulations as to the real data. For the NFW profile (or closely related profiles), this has recently been done by a number of groups \\citep{bek11,ogh11,bmk11}. The intrinsic scatter due to projection effects is found to be $\\sim 25$\\% \\citep{bek11,bmk11}, while the expected bias is dependent on the outer fit radius -- if this is restricted to be close to the virial radius, the average mass can be recovered with little bias \\citep{bek11,ogh11}.  The unbiased mean, yet considerable intrinsic scatter, for cluster weak lensing measurements implies that relatively large samples of clusters are necessary to meet the calibration needs of cluster cosmology. For such work, the clusters used should ideally be drawn representatively from the surveys in question, so as to have the same selection function. This is a fundamental reason why, for example, strong-lensing selected clusters should not be used for this purpose -- the incidence of strong lensing is highly biased towards clusters that are elongated and/or have additional structures along the line of sight.  To date, only a handful of studies have measured individual weak-lensing masses for more than a few clusters, and none have fully incorporated the results into a robust cosmological work, which would require solving simultaneously for the scaling relations and cosmological parameters \\citep{mar10,aem11}.  A number of early works \\citep{all98,hok98,csk04,sks05} compared lensing mass estimates of massive clusters to X-ray mass proxies, but the weak-lensing mass measurements were generally limited by the small fields of view of existing cameras. The work of \\citet[][see also \\citealt{dah06,ped07}]{dki02} provides the so-far largest compilation of weak-lensing mass measurements of individual clusters (38 clusters).  With the increasing availability of high-quality, wide-field mosaic cameras, the precision of weak-lensing mass measurements at sufficiently large cluster radii has significantly increased, providing the means to study cluster scaling relations with total mass measurements.  \\citet{hoe07} compared weak-lensing masses of 20 clusters, derived from two-filter optical imaging, to independently measured X-ray luminosities and temperatures, as well as galaxy velocity dispersions. For 18 of these clusters, \\citet{mhb08} computed X-ray hydrostatic masses and compared these to the weak-lensing mass estimates.  \\citet{bsk07} compared weak-lensing masses for 11 clusters measured from three-filter imaging to X-ray luminosities and temperatures.  The LoCuSS project measured weak-lensing masses with two-filter imaging for 30 clusters \\citep{otu10}.  For 12 of them, the lensing masses were compared to $T_X$, $M_{\\rm gas}$, and $Y_{\\rm X}$ \\citep{ozf10} as well as hydrostatic mass estimates \\citep{zof10}. Using 18 of these clusters, \\citet{mso11} present a first comparison of integrated Compton parameters from SZ observations to weak-lensing mass determinations. \\citet{hhl12} present a second SZ--weak lensing comparison for 5 clusters. \\citet{hdc11} used single-filter Hubble Space Telescope observations to measure weak-lensing masses for 25 clusters of moderate X-ray luminosity, and compared these to the cluster X-ray luminosities and temperatures. For larger samples of less massive systems, stacking analyses enable the determination of the mean cluster mass in bins of survey observable \\citep{jsw07,lfk10}. For most studies listed here, the bulk of the cluster samples studied is at $z_{\\rm Cl} \\sim 0.2 - 0.3$. A few studies have specifically targeted higher-redshift clusters, both using space-based \\citep[][22 clusters at $z\\gtrsim1$]{jdh11} and ground-based imaging \\citep[][7 clusters at $z \\sim 0.4 - 0.8$]{ier11}.  A key assumption of these pathfinding lensing studies is to implicitly place all background galaxies at the same effective redshift.  For low-redshift clusters ($z_{\\rm Cl}\\sim0.2$, representing the bulk of the clusters studied to date), this approximation should not severely bias the mass measurements; the peak of the galaxy distribution is at $z\\sim0.8-1.0$ and, for clusters at low redshifts, the shear signal varies only slowly over this range, causing errors in the effective redshift to bias the mass only slightly.  For clusters at higher redshifts ($z_{\\rm Cl}\\gtrsim0.4$), however, this is no longer the case and one can significantly reduce systematic scatter and potential bias (in case the assumed redshift distribution is not representative of the redshift distribution in cluster fields) by incorporating appropriate redshift information for individual galaxies. Since weak lensing is based on shape measurements of many faint galaxies, this is feasible only with photometric redshifts.  Current and up-coming cluster surveys, such as the South Pole Telescope survey \\citep[SPT, ][]{vcd10}, the Atacama Cosmology Telescope \\citep[ACT, ][]{sta11}, Planck \\citep{pla10}, the Dark Energy Survey \\citep[][DES]{des05}, and eROSITA \\citep{pab10} will find hundreds to many thousands of massive clusters in the redshift range $0.5 \\lesssim z \\lesssim 1.5$. Lensing mass calibrations for these surveys will be vital to maximizing their potential to constrain cosmology. It is therefore essential to develop the strategies and tools to measure unbiased cluster masses using photometric redshifts in an optimal way.   \\subsection{This study}  In this series of papers, we develop and apply techniques to enable the determination of accurate weak-lensing masses for a total of \\clusterfields\\, clusters from deep, high-quality multi-color Subaru SuprimeCam and CFHT MegaPrime optical imaging.  In this first paper, we describe the cluster sample and the data reduction methods: a careful data treatment is key to robust shear and photometry measurements, and unbiased cluster mass determination. We discuss the correspondence between the dark matter, gas, and optical light distributions, and the relation of the positions of the Brightest Cluster Galaxies (BCGs) and X-ray centroids. In Paper II \\citep{kla12}, we describe the details of our photometric calibration, including a prescription to construct the ``star flat'', which corrects flat-field errors due to varying pixel scale and scattered light in wide-field cameras. In Paper II we also describe an improved and versatile technique to calibrate photometric zeropoints from stellar colors, whose implementation we have made publicly available. Using these methods, we show that we can estimate robust photometric redshifts even when calibration data are lacking, and present an initial analysis of the source-redshift dependent shear signal of the clusters. In Paper III \\citep{alk12} we introduce a novel Bayesian approach to weak-lensing mass estimation that makes full use of photometric redshift probability distributions of lensed galaxies. We compare the obtained masses to those derived from the more common method of adopting a single effective redshift for the background galaxies. Critically, we also include a detailed discussion and quantification of the systematic uncertainties involved. Additional papers will focus on the scaling relations between weak-lensing masses and other observables, and present updated cosmological constraints.  For a project such as this, where a central goal is the comparison of measurements determined by independent techniques, and where the measurements to be calibrated have already been used in cosmological studies, there is a clear danger of ``observer bias'' or ``confirmation bias''. These biases are well-known in the wider physics community, and can be avoided by implementing ``blind analyses'' \\citep{klr05}. While blind analyses have not yet been used widely in astronomy to date, they will be essential for upcoming precision cosmology measurements \\citep[see also][]{aem11,crd11}. To combat confirmation bias, we have chosen to explicitly avoid direct comparison with X-ray mass proxies, or indeed any other mass estimates, during the course of this study, revealing all comparisons only at the end of a given part of the study.  To enforce this restriction, in the few cases where intermediate results were presented, all non-lensing mass estimates were multiplied by a random, unrevealed number (all masses were multiplied by the same number), thus removing the absolute scalings of the lensing vs. other mass relations, the primary quantities of interest. Since the lensing data are not altered, this procedure allows complete and accurate analyses of statistical and systematic errors, while eliminating unintentional bias towards the expected correlation with other mass proxies.  In the very early stages of this work blinding was not implemented and preliminary comparisons of crude mass estimates for a small fraction ($\\lesssim 20$ per cent) of the clusters were examined. We emphasize, however, that the final, more sophisticated mass measurement methods described in Paper III were developed independently of these early analyses and that all lensing mass measurements presented in these papers were determined blindly with respect to other mass proxies and all results in the literature.  \"Unblinding\" with respect to lensing mass estimates in the literature took place only after the lensing analysis was completed, including internal review of Papers I-III. \"Unblinding\" with respect to X-ray and other independent mass proxies has not occurred at the time of completing papers I-III. Any subsequent changes to the lensing analysis will be reported in Paper III, or later work, if necessary.  This paper is structured as follows: In Sect.~\\ref{sect:clusters} we describe the cluster sample and the optical imaging observations. Since the lensing analysis is performed mostly on SuprimeCam data, we give a detailed description of the SuprimeCam data reduction in Sect.~\\ref{sect:optical_data} (with additional details in App.~\\ref{appendix:early_data}). Sect.~\\ref{sect:catalogs_photom} describes the object detection and initial photometry measurements. In Sect.~\\ref{sect:shearmeasurements} we briefly summarize the shear measurement method based on \\citet{ksb95}, discuss our strategies to correct for the anisotropy (also App.~\\ref{appendix:psf}) and isotropic smearing of the point spread function, and discuss the calibration using STEP2 simulations \\citep{mhb07}, including accounting for correlated noise. In Sect.~\\ref{sect:cluster_maps} and App.~\\ref{appendix:clustermaps} we present a gallery of cluster images and maps of the total mass distribution as recovered from the weak-lensing data, the large-scale structure around each cluster as traced by galaxies on the red sequence, and the X-ray emission associated with the cluster.  In Sect.~\\ref{sect:cluster_centers}, we investigate the impact of different choices for the cluster centers on the lensing results. We summarize and provide an outlook on future work in Sect.~\\ref{sect:summary}.  The fiducial cosmology adopted in this paper is a flat $\\Lambda$CDM model with $\\Omega_{\\rm m} = 0.3$ and $H_0 = 100 \\,h\\, \\mbox{km/s/Mpc}$, where $h=0.7$. ",
        "conclusions": "Here we provide a brief summary of the KSB algorithm. For a more in-depth review, see e.g. \\cite{bas01}.  In the KSB algorithm, the complex ellipticity $\\bm{e}$ for each object is estimated from the second moments $Q_{ij}$ of the object's light distribution $I(\\vec{\\theta})$: \\be e = e_1 + ie_2 = \\frac{Q_{11} - Q_{22} + 2 i Q_{12}} {Q_{11} + Q_{22}} \\quad , \\ee where \\be Q_{ij} = \\int \\dif^2 \\theta\\: I(\\vec{\\theta}) \\:  W_{r_g}(|\\vec{\\theta}|)  \\: \\theta_i \\theta_j \\quad , \\label{eq:2nd-moments} \\ee and $W_{r_g}$ is a Gaussian weight function of width $r_g$ \\citep[we use the {\\tt FLUX\\_RADIUS} measured by {\\sc SExtractor} for $r_g$; see the discussion in][]{sch08}.  Following \\citet{evb01}, we use the same weight function to define a signal-to-noise ratio of each object: \\be S/N = \\frac{\\int \\dif^2 \\theta I(\\vec{\\theta}) W_{r_g}(|\\vec{\\theta}|)} {\\sigma_{\\rm sky} \\sqrt{\\int \\dif^2 \\theta W_{r_g}^2(|\\vec{\\theta}|) }} \\quad , \\label{eq:lensing-snr} \\ee which captures the uncertainty in the shape measurement. This shape-specific signal-to-noise measure is different from the flux signal-to-noise: for our data, objects with a shape $S/N$ of $\\sim 3$ are highly significant detections with ${\\tt FLUX\\_ISO} / \\sigma_{\\tt FLUX\\_ISO} \\sim 30$.   Gravitational lensing has two effects on the observed shapes of background galaxies: the convergence $\\kappa$ scales the image of a background object isotropically, and the shear $\\bfmath{\\gamma}$ stretches it anisotropically.  The combined effect is the reduced shear: \\be \\bm{g} \\;=\\; \\frac{\\bfmath{\\gamma}}{1-\\kappa}\\;. \\ee  In the limit of small shear ($g\\ll1$), and only small anisotropy of the telescope's point spread function (PSF), the (seeing-convolved) intrinsic ellipticity $\\bm{\\hat{e}}^0$ of an object is transformed to the observed ellipticity \\be \\bm{e} \\;=\\; \\bm{\\hat{e}}^0 +  P^{\\rm g}\\bm{g} + P^{\\rm sm}\\bm{q}^{\\star}\\; ; \\; P^{\\rm g} = P^{\\rm sh} - P^{\\rm sm}(P^{\\star \\rm sm})^{-1} P^{\\star \\rm sh}\\;. \\label{eq:ksb} \\ee The stellar anisotropy kernel $\\bm{q}^{\\star}$ describes the anisotropic component of the PSF; the smear polarizability tensor $P^{\\rm sm}$ describes the susceptibility of an object to the PSF anisotropy (and largely depends on the apparent object size); and the shear polarizability tensor $P^{\\rm sh}$ describes the object response to the shear. $P^{\\rm sm}$ and $P^{\\rm sh}$ are measured from an object's third and fourth order moments. The starred quantities of these tensors are measured on stars; but note that the weight function must be adjusted to the object size \\citep{hfk98}.  $\\bm{q}^{\\star}$ is measured from stars (for which the gravitational shear $\\bm{g}$ and intrinsic ellipticity $\\bm{\\hat{e}}^0$ vanish), so that the galaxy ellipticities can be corrected for the anisotropy of the PSF. The reduced shear is then \\be \\bm{g} = (P^{\\rm g})^{-1} (\\bm{e}^{\\rm aniso} - \\bm{e}^{0})\\quad ;  \\quad \\bm{e}^{\\rm aniso} = \\bm{e} - P^{\\rm sm}\\bm{q}^{\\star} \\quad . \\ee The source ellipticity $\\bm{e}^{0}$ is of course not known. KSB instead returns \\be \\hat{\\bm{g}} = (P^{\\rm g})^{-1} \\bm{e}^{\\rm aniso} \\quad . \\ee  Because galaxies are randomly oriented (at least to the precision required for cluster weak lensing), the average ellipticity of an unlensed population of galaxies vanishes: $\\ave{\\bm{e}^{0}} = 0 = \\ave{(P^{\\rm g})^{-1}\\bm{e}^{0}}$. Hence, the justification for KSB is that the expectation value $\\ave{\\hat{\\bm{g}}}$ is an estimate of $\\bm{g}$.  Since the trace-free part of the $P^{\\rm g}$ tensor is much smaller than the trace, we follow \\citet{evb01} and make the approximations \\be (P^{\\star \\rm sm})^{-1} P^{\\star \\rm sh} \\rightarrow \\frac{{\\rm Tr}[P^{\\star \\rm sh}]}{{\\rm Tr}[P^{\\star \\rm sm}]} =: T^{\\star} \\quad ; \\quad (P^{\\rm g})^{-1} \\rightarrow \\frac{2}{{\\rm Tr}[P^{\\rm g}]} \\ee which also reduces sensitivity to noise \\citep{hvb06}.   \\subsection{Shape measurements and star selection} \\label{sect:shapemeassubsect}  \\begin{figure*} \\includegraphics[width=0.49\\hsize]{figures/MACS0025-12_rh_mag_pretty_lores}\\hspace{0.01\\hsize}% \\includegraphics[width=0.49\\hsize]{figures/MACS0025-12_psfcorr_pretty} \\caption{Illustration of the star selection and PSF anisotropy correction, shown for the \\bbullet$\\;$ {\\it V}$_{\\rm J}$ field. Panel (a) shows the magnitude - radius diagram for all objects in the catalog. The stellar sequence is clearly visible at $r_h \\sim 1.4$px. The preselection of stars is shown as the red box. Stars that are not rejected as significant outliers in the initial second-order polynomial fit are shown in green. In the figures on the right, we illustrate the correction itself. Panel (b) shows the uncorrected stellar ellipticity pattern: at the position of each star, the measured ellipticity is indicated as a line with a length proportional to $\\left|e^{\\star}\\right|$, with orientation $\\phi = 0.5 \\arctan(e_2^{\\star}/e_1^{\\star})$ . Panel (c) shows the distribution of $e_1^{\\star}$ vs. $e_2^{\\star}$ values. Panel (d) shows the residual ellipticities after correcting the PSF pattern with an eighth-order polynomial (see Fig.~\\ref{fig:psf2}); the distribution of corrected $e_1^{\\star}$ and $e_2^{\\star}$ values is shown in panel (e). Note that for the SuprimeCam data, the size of the coadded output image ($33{\\rm arcmin}\\times33{\\rm arcmin}$) is larger than the area covered by the input frames -- because we mask pixels more than 15~arcmin from each exposure center, the non-zero-weight part of the output image appears roughly circular. The ``missing corners'' visible in panels (b) and (d) are due to the top-left chip ``w67c1'', which has noticeable CTI and lower QE than the other chips, and therefore is rejected.} \\label{fig:psf} \\end{figure*}  For the shape measurements, we first run {\\sc SExtractor} in dual-image mode, with the median coadded image as the detection image, and the image coadded for lensing (coadded with a weighted average) as the second, ``measurement'', image. This step mainly serves the purpose of obtaining the {\\sc SExtractor} {\\tt FLUX\\_RADIUS}($=r_g$) measurement for each object on the actual lensing image, while retaining the same object selection and identification as for the photometric catalogs.  For objects with $0.5 \\le r_g \\le 10$ we measure the ellipticities as described above with the code {\\sc analyseldac} \\citep{evb01}.  Larger objects are unlikely to be background galaxies; while smaller objects are predominantly spurious detections.  {\\sc analyseldac} also provides a more robust measure of the half-light radius, $r_h$. We select stars for the PSF correction in a diagram of magnitude vs. $r_h$, where stars with sufficient signal-to-noise, but which are not saturated, form a well-defined sequence (Fig.~\\ref{fig:psf}). The star selection is refined by fitting the PSF anisotropy across the field with a second-order polynomial, and rejecting 5$\\sigma$ outliers (for the actual PSF correction we use a higher-order polynomial; see next section). The number of stars varies considerably in the sample, from 300 to 3000 per field. With the roughly circular field of view of radius $\\sim 15$~arcmin, this corresponds to $0.4 - 4 {\\rm \\; stars \\; arcmin}^{-2}$, with the typical number density being $\\sim 1 {\\rm \\; star \\; arcmin}^{-2}$. The stars selected here also form the basis of the relative photometric calibration between bands via the stellar color-color locus (see Paper II); the color-color diagrams confirm that the stellar sample selected here is fairly clean.   \\subsection{PSF anisotropy correction} \\label{sect:psfcorr}  Correcting for the PSF anisotropy, $\\bm{q}^{\\star} = (P^{\\star \\rm sm})^{-1} \\bm{e}^{\\star}$, is essential, as it can mimick shear. However, the PSF can only be measured at discrete locations in the image plane, namely at the positions of suitable stars.  We measure $\\bm{q}^{\\star}$ for the selected stars (see above), and fit a polynomial function to the spatial variation in both components (Fig.~\\ref{fig:psf}).  An important question here is whether the PSF is stable across chip boundaries -- this is the case if the CCDs are sufficiently coplanar. If the CCDs are mounted at different heights, the focus position and hence PSF shape changes abruptly across the CCD boundary. On a single exposure, or coadded images with dither patterns of the size of the gaps between chips, this could be accounted for by fitting the PSF variation for each CCD separately. For our data, however, the dither patterns are significantly larger, and for most fields, the camera has been rotated by 90$^{\\circ}$ between exposures. We have tested for ``jumps'' of the SuprimeCam PSF in images with excellent seeing and a large density of stars (see App.~\\ref{sect:planarity}) and find that in configurations 10\\_1 and 10\\_2 (used for shape measurements), the CCDs are remarkably coplanar -- there are no measurable discrete PSF jumps across chip boundaries. (This is not the case in early data, another reason to disregard those data for lensing purposes.) Hence, the PSF pattern can be corrected across the full field of view, without the need to correct on a chip-by-chip basis.    \\begin{figure*} \\includegraphics[trim=0.3cm 1.3cm 0.5cm 2.5cm,width=0.52\\hsize]{figures/MACS0025-12_W-J-V_good_PshstPsmst_9rm_idl}\\hspace{0.03\\hsize}% \\includegraphics[width=0.45\\hsize]{figures/PshstPsmst_pretty} \\caption{Illustration of the PSF isotropy correction as a function of position and object size, for the \\bbullet$\\;$ {\\it V}$_{\\rm J}$ field.  The left panel shows the variation of $T^{\\star} = {\\rm Tr}[P^{\\star \\rm sh}]/{\\rm Tr}[P^{\\star \\rm sm}]$ across the field of view. At the position of each star, we indicate $T^{\\star}$, measured with a weight function of width $r_g=0.6$arcsec. Note the 20\\% variation across the image, which we fit with a second-order polynomial. The right panel shows $T^{\\star}$ as a function of object size, evaluated at each object position with the appropriate weight function. The spread in $T^{\\star}$ at a given $r_g$ reflects the spatial variation shown on the left. We impose a minimum size criterion of $r_g>1.5 {\\rm px}$ (illustrated by the dotted blue line), to avoid the upturn at smaller scales. For comparison, the median $r_g$ of stars is indicated by the dashed blue line. } \\label{fig:isocorr} \\end{figure*}   The PSF of SuprimeCam (and MegaPrime) can vary considerably over the field of view, even in single exposures. We fit the entire field with a single polynomial, but find that usually a high order polynomial is required (from fourth order up to a limit of tenth order). Other authors instead divide the field into subsets and fit these with second order polynomials, but since this creates discontinuities in the PSF, we prefer the single, higher-order polynomial.   We developed a number of criteria to judge the quality of the PSF correction, and to choose the minimum polynomial order required to achieve a good fit. This process is described in App.~\\ref{sect:psfquality}.  Note that we calculate $\\bm{q}^{\\star}$ using the weight function of each star (i.e. $r_g = r_g^{\\star}$). As \\citet{hfk98} have argued, all quantities in Eq.~\\ref{eq:ksb} should be measured with the same weight function as the object (galaxy) to be corrected. However, if the anisotropy of the PSF does not vary with isophote level (which is a good approximation for many ground-based instruments), $\\bm{q}^{\\star}$ is independent of the width of the weight function. Measuring it with the stellar weight function automatically reduces the noise in this measurement, making the PSF measurement more robust. We find no systematic shift in shear measurements when measuring $\\bm{q}^{\\star}$ with the galaxy weight function, consistent with the results of \\citet{hvb06}.   \\subsection{PSF isotropy correction} \\label{sect:psfisocorr}   The isotropic part of the PSF (expressed as the $P^{\\rm g}$ tensor), circularizes object shapes. If inadequately corrected, this can lead to a dilution of the shear measurement. For the calculation of $P^{\\rm g}$, $T^{\\star} = {\\rm Tr}[P^{\\star \\rm sh}]/{\\rm Tr}[P^{\\star \\rm sm}]$ needs to be measured from stars. This quantity is sensitive to the size of the weight function and therefore must be measured with the weight function appropriate for the object to be corrected, and within the same aperture used for the object. Furthermore, $T^{\\star}$ can vary spatially, as the size of the PSF can vary within the field of view. In well-focused exposures, the PSF tends to be smaller at the center of the field of view than towards the edges. If this is left unaccounted for, it can lead to systematic biases in the radial shear profile, and thus the cluster mass measurement. We measure $T^{\\star}$ at discrete values $r_g^b$ of the weight function size over the range $0.33 \\le r_g^b \\le 10$, in 0.33 pixel increments. For each weight function scale, we fit the spatial variation of $T^{\\star}$ with a second-order polynomial across the images, which suffices to capture the variation (Fig.~\\ref{fig:isocorr}).  For each object, we then assign $T^{\\star}$ according to the fit for $r_g^b$ closest to the object size $r_g$. The trend of $T^{\\star}$ with object size is linear for objects larger than the PSF, but shows an upturn at $r_g \\lesssim 1.5$~px (right panel of Fig.~\\ref{fig:isocorr}). Since this upturn is likely an artefact, we reject objects with $r_g<1.5$~px. (Note that in the lensing analysis, we apply an additional, stricter size criterion based on $r_h$; see Paper~III.)   \\subsection{Coaddition and PSF correction -- influence on cluster mass measurements?} \\label{sect:psf_influence_on_mass}  Apart from spatial variation, the SuprimeCam PSF is also temporally variable.  Both the telescope and camera contribute to the anisotropy of the PSF; i.e. in some fields, rotating the camera by $90^{\\circ}$ causes stellar ellipticities to reverse sign (in sky coordinates), and sometimes the PSF pattern remains largely intact through rotation. Since for a significant fraction of our data, the camera was rotated between exposures, we must ask whether the PSF of a coadded image can still be adequately corrected. The goal of our project is to measure unbiased cluster masses, and so we evaluate this issue by testing whether masses measured from coadded images are biased. We perform this test on the ``worst-case'' fields, where several sets with very different PSF patterns have been coadded. These are fields with excellent (but still adequately sampled) seeing ($\\sim 0.5\\arcsec$), with camera rotation between exposures, and with exposures often taken on more than one night (note that on all but one field, the size of the seeing disk is comparable between nights due to our lensing image selection process). Within a given field, night, and rotation, the PSF pattern is relatively stable. We therefore coadd images from these subsets (i.e. for a given subset, only exposures from the same night and with the same camera rotation angle are used), and compare masses measured on these images to the image coadded from all subsets (the mass measurement is described in Paper III).  \\begin{figure} \\includegraphics[width=\\hsize]{figures/compare_masses_coadds} \\caption{The distribution of ratios between masses measured from lensing images coadded from single nights and rotations and masses measured from the full coadded lensing image (using all available nights and rotations). The masses were measured with the ``color-cut'' method; see Paper III. No bias is introduced into the mass measurements by using full coadded lensing images. The scatter in mass ratios is largely due to limited number statistics in some images; e.g. the ``outliers'' with values of 0.5 and 1.5 were measured from less than 2000 galaxies.} \\label{fig:compare_masses_coadds} \\end{figure}  The distribution of ratios between masses measured on subsets and on the fully coadded image is shown in Fig.~\\ref{fig:compare_masses_coadds}.  For most comparisons, mass measurements from different coadded images agree very well: the distribution clearly peaks at 1, with a median ratio of 0.993, indicating that using the full coadded image does not lead to biased cluster mass measurements on average.  We furthermore investigate the influence of the PSF correction on the mass measurements. For this purpose, we compare masses determined with shear measurements corrected only with a second order polynomial to those using the polynomial order determined according to the process described in App.~\\ref{sect:psfquality} (almost all sixth, eighth, or tenth order). The change in measured mass when the PSF is not adequately corrected is best described by an offset of $\\sim - 5 \\times 10^{13}M_{\\odot}$, measured at $2.5\\sigma$ significance. For clusters in the mass range of our sample, this corresponds to mass underestimates of the order of 1--10\\%, illustrating the requirement of an adequate PSF correction for cluster mass measurements.  To test whether the mass measurements are robust against changes in the details of how the PSF correction polynomial is determined, we also compare mass measurements if the order of the polynomial is decreased or increased by two orders. There is no significant mass shift. The PSF correction criteria developed in App.~\\ref{sect:psfquality} therefore are sufficient for our purpose.    \\subsection{STEP calibration} \\label{sect:step}  \\begin{figure*} \\includegraphics[width=0.48\\hsize]{figures/step_snratio} \\includegraphics[width=0.48\\hsize]{figures/step_size} \\caption{Results of our calibration of the shear measurement bias from the STEP simulations. Shown is the multiplicative shear bias $m$ (i.e. a value of $-0.1$ means the estimated shear is 90\\% of the true shear, since the additive bias is small and consistent with zero), as a function of $S/N$ (left panel) and object size (right panel). For $(S/N)_{\\rm STEP}\\gtrsim 7$, the bias is approximately constant, $m \\sim -0.09$ for PSF A (0.6\\arcsec), and $m \\sim -0.06$ for PSF C (0.8\\arcsecf). Below this threshold, the magnitude of the bias is significantly larger. However, because of the strong correlated noise in the STEP2 images, $(S/N)_{\\rm STEP}$ overestimates the true signal-to-noise ratio. Our images are less susceptible to correlated noise (due to choosing the {\\sc Lanczos3} kernel for resampling; see text for details). We find that the $(S/N)_{\\rm STEP}\\gtrsim 7$ threshold approximately corresponds to $(S/N)_{\\tt Lanczos3}\\geq3$ for our images, and hence impose this criterion for objects entering the shear analysis. The right panel shows the shear bias for objects with $(S/N)_{\\rm STEP}\\gtrsim 7$ as a function of object size, in units of the PSF size (measured as the median $r_h$ of stars selected for the PSF correction). There is a notable trend with size, in that the magnitude of the bias is larger for smaller objects. To correct for this size-dependence, we fit a piecewise linear function to the unbinned data, constrained to be a constant value for large objects. We find no statistically significant difference between corrections for the two shear components.  The dotted line indicates the minimum size criterion used to reject point sources.} \\label{fig:shapecorr} \\end{figure*}  A crucial element of the analysis is to calibrate the shear measurement bias inherent to KSB methods, which typically underestimate the shear \\citep{evb01}. Any underestimate of the shear will result in a direct underestimate of the cluster mass. Fortunately, the cosmic shear community has led efforts to provide calibration datasets for shear measurement methods with the Shear TEsting Programme \\citep[STEP,][]{hvb06,mhb07}. We use the simulations from the STEP 2 project \\citep{mhb07} to calibrate the estimator $\\hat{\\bm g}$ to the true input shear ${\\bm g}$ as a function of the $S/N$ and size of each galaxy. In Paper III we describe the characterization of the full probability distribution of $p(\\hat{\\bm g} | \\bm g)$; here we summarize the results applicable when considering simple averaging of the shear estimators.  \\subsubsection{Shear measurement bias as a function of S/N}  Fig.~\\ref{fig:shapecorr} shows the results for the average multiplicative shear bias, determined from the STEP data, as a function of the signal-to-noise ratio measured in the STEP images, $(S/N)_{\\rm STEP}$. The bias is highly dependent on $(S/N)_{\\rm STEP}$, in the sense that it is consistent with a constant value above $(S/N)_{\\rm STEP}\\gtrsim7$, and increases significantly (in magnitude) for objects with lower $(S/N)_{\\rm STEP}$. We find a slightly smaller correction for PSF model ``C'' with 0.8\\arcsec\\, seeing, than for PSF ``A'' with 0.6\\arcsec, possibly because the PSF is better sampled.  Of the PSF models tested in \\citet{mhb07}, these are the most appropriate; since we discard the edges of the field, the highly elliptical PSF ``D'' and ``E'' are less applicable.  This behavior, with an approximately constant shear bias above a threshold signal-to-noise, and large negative bias below the threshold, is very similar to other KSB implementations \\citep{mhb07}. In particular, it is very similar to the ``TS'' implementation of \\citet{ses07} tested in STEP-2. In \\citet{hss09}, those authors also show the shear bias as a function of $S/N$. Taking into account the re-scaling of $S/N$ due to correlated noise (see below) and the fact that TS apply a constant shear scaling of 1.08, the results are in excellent agreement \\citep[as would be expected, as both methods are based on the same implementation of KSB+ as described in][]{evb01}. Using all the PSF models tested in STEP-2, these authors detect and subsequently correct for a slight $S/N$-dependence also for large $S/N$ \\citep{shj10}. In the two PSF models we tested, we do not find evidence for such a trend with $S/N$, but acknowledge that we might be lacking enough statistics to do so.  Fig.~\\ref{fig:shapecorr} illustrates strikingly the need for accurate $S/N$ estimates for each object. By requiring $(S/N)_{\\rm STEP}>7$, we can ensure that we utilize only objects in the regime where the shear measurement bias is robust, and does not depend sensitively on the signal-to-noise ratio.   \\subsubsection{Accounting for correlated noise} \\label{sect:corr_noise}  A complication with applying the STEP calibration to the actual images is that the measured noise properties (and thus $S/N$) are sensitive to correlated noise. Signal-to-noise in images with correlated noise is overestimated if the estimation procedure does not explicitly account for its presence. Corrections to account for the effects of correlated noise have been made to address correlated noise in photometry measurements \\citep{cmd00,masci09} and in shape measurements \\citep{ses07}.  In the actual data, the correlated noise stems from the resampling process, but the choice of the {\\tt Lanczos3} kernel minimizes the amount of correlation introduced. In the STEP2 images, correlated noise was artifically introduced by smoothing with a gaussian kernel.  The effects of correlated noise are most pronounced when only a few pixels are involved in a measurement, and is asymptotic to a constant correction at a large number of pixels. To quantify the effects of correlated noise on $S/N$, we created artificial images with Gaussian noise approximating the noise properties of our images, and 10000 ``objects'' (for simplicity, we use normal distributions of a fixed width) at equal spacing in the images. We then resampled these images using the same kernel as our actual images, as well as applying a Gaussian smoothing kernel, as was done for the STEP2 images.  By comparing the distribution of measured $S/N$ values for the shape measurement procedure for the two resampled images, we are able to determine the scaling factor between the signal to noise measure on our images and the STEP2 images. To test for aperture size effects, we repeat this procedure for objects with FWHM of 2 pixels to 18 pixels. There is only a weak dependence of the scaling on galaxy size over the range of galaxies accepted into our analysis. We estimate the scaling of $S/N$ between our images to the STEP images to be $\\approx2.3$. The threshold in shear bias that we see in the STEP images at $(S/N)_{\\rm STEP}\\sim7$ thus corresponds to $S/N\\sim3$ for our data. By requiring $S/N\\ge3$, we therefore robustly select only objects for which the average shear calibration bias is not a strong function of $S/N$.  The significance $\\nu$ that is used in the {\\sc imcat} implementation of KSB is larger than the $S/N$ measure by a factor of 3.5 \\citep{evb01}; our $S/N$ cut therefore corresponds to $\\nu \\gtrsim 11.5$. Compared to other weak-lensing studies, this $S/N$-cut is relatively conservative \\citep[see e.g. Table A1 of][]{hvb06}. It is also worth keeping in mind that the lensing-$S/N$ is lower than the detection significance (see Sect.~\\ref{sect:ksbsummary}).     \\subsubsection{Assigning the correct S/N ratio}  A further practical consideration is necessary to assign the correct $S/N$ for each object, because {\\sc analyseldac} assumes a constant sky noise level when calculating $S/N$ (Eq.~\\ref{eq:lensing-snr}). For our images, with large dither patterns, rotation between exposures, and ample masking, this is certainly not the case, as illustrated by a typical weight map shown in Fig.~\\ref{fig:weightmask}.  We correct for this by scaling the reported $S/N$ according to the local weight. The sky noise used to calculate $S/N$ from Eq.~\\ref{eq:lensing-snr} is the average RMS of the sky background as measured by {\\sc SExtractor} on the lensing image. The relative sky noise scales with exposure time as $\\sigma_{\\rm sky}\\propto 1/\\sqrt{t_{\\rm exp}}$ (recall that our images are normalized to a 1s exposure time). The weight map tracks the effective exposure time per pixel; therefore we can recover the actual $S/N$ by scaling the reported value by the square root of the ratio between the average non-zero weight and the local weight. By comparing with measurements made on smaller image cut-outs with constant noise, we have verified that this recovers the true $S/N$ to within a few percent, enough precision to identify galaxies above the threshold $S/N$ value.    \\subsubsection{Shear measurement bias as a function of object size}  In the regime where the average shear measurement bias does not depend on $(S/N)_{\\rm STEP}$, we test for dependence on the size of the object (in units of PSF size, Fig.~\\ref{fig:shapecorr}). We find that the shear underestimate is worst for objects just larger than the PSF, and is smallest for well-sampled objects. This is expected and consistent with other KSB implementations \\citep{mhb07}. We therefore express the correction to be applied to the shear measurement as a function of object size ($r_h$ as returned by {\\sc analyseldac}, scaled by the size of the PSF, defined as the median $r_h$ of stars selected for the PSF correction). Fig.~\\ref{fig:shapecorr} illustrates the best-fit correction; in Paper III this process is described in more detail, including how the uncertainties in the shape correction are propagated to the mass measurements.    \\label{sect:summary}  This is the first of a series of papers aimed at measuring accurate weak-lensing masses for \\clusterfields\\, of the most X-ray-luminous galaxy clusters -- the true giants in the observable Universe. The primary goal is to measure the key mass--observable scaling relations for clusters to better than 10\\% accuracy, a vital prerequisite for current and future cluster surveys to utilize their full statistical power. To achieve this goal, we have developed new methods and improved upon existing ones to measure accurate weak-lensing cluster masses, and have rigorously quantified the residual sources of systematic uncertainty.  The cluster sample presented here is the largest to-date for which weak-lensing masses have been measured with a homogeneous dataset and methodology. With a redshift range of $0.15 \\lesssim z \\lesssim 0.7$, and with half the clusters at $z>0.4$, it extends to higher redshifts than previous ground-based studies. However, its key distinction is the emphasis on the minimization and accurate quantification of residual systematic uncertainties, and the blind nature of the lensing mass analysis with respect to other mass proxies.  Although the intrinsic scatter of 3D weak-lensing mass measurements is large ($\\sim\\!30$\\%), cluster sample sizes of $\\sim\\!50$ bring the statistical uncertainty on the mean cluster mass (or mean ratio of weak-lensing mass to other mass proxy) to the 5\\% level. Hence, systematic uncertainties should ideally be controlled to the level of a few percent in order not to limit weak-lensing mass calibration efforts.  There are three main sources of systematic uncertainties for 3D mass measurements of individual clusters: the two observational challenges lie in measuring unbiased estimators of galaxy shapes, and their redshifts. The third source of systematic uncertainty lies in relating the measured shear and redshift estimates to the mass of the cluster.  In this paper, we have laid the basis for the subsequent lensing analysis, describing a robust data reduction method that aims for both excellent shape measurements and photometry measurements.  We show that the shear bias of the KSB method is a strong function of signal-to-noise ratio, and a function of object size even after low signal-to-noise objects have been rejected.  Assigning each object the appropriate shear calibration is critical to this work, in particular when individual photometric redshift estimates are used. We address one source of uncertainty for the relation between measured shear and cluster mass, namely the choice of cluster center. We adopt the centroid of the X-ray emission as the cluster center, and show that the location of the dominant cluster galaxy, if correctly identified, agrees well with the X-ray centroid, with a median projected offset of only 20~kpc.  Only for the most extreme bimodal cluster mergers, such as \\mos and \\bbullet, are the BCG and the X-ray centroid separated by $\\gtrsim$~100~kpc. We find no systematic bias between weak-lensing mass measurements centered on the BCGs compared to those relative to the X-ray centroids. For the clusters considered here, and with our weak-lensing methodology, miscentering therefore is not a source of systematic uncertainty.  For each cluster, we show optical images and maps of the total mass distribution measured from weak lensing, cluster structure and surrounding large-scale structure as traced by red-sequence galaxies, and the extended X-ray emission. These multi-wavelength maps illustrate the large-scale structure within which each cluster is embedded, as well as possible interactions with other mass concentrations in the field and the presence of foreground and background structures.  In Paper~II \\citep{kla12}, we detail the key methods used for accurate photometric calibration and determination of photometric redshifts. We describe how to correct position-dependent flux-zero-points from repeated observations of the same fields (and SDSS photometry, when available).  We describe a number of improvements to the ``stellar locus method'', allowing us to precisely calibrate the relative zero-points between filters through reference to the narrow intrinsic locus of main sequence stars in color-color space. We show that with these techniques, we recover robust photometric redshifts even in the absence of calibration data. The quality of the photometric redshifts is further illustrated with the shear-redshift scaling behind the clusters.  In Paper~III \\citep{alk12}, we proceed to the actual mass measurements. We develop a novel Bayesian algorithm which utilizes the full photometric redshift probability distribution, and the full distribution of KSB shear estimates with respect to the true shear. We extensively test this method on the COSMOS field and show that, in terms of the mean recovered cluster mass, the bias of our method is at most 2\\% over the cluster redshift range considered here. We also measure cluster masses using an improved version of the traditional ``color-cut method'' used in other works, which is applicable to a larger number of weak-lensing cluster datasets.  We make detailed estimates of the residual systematic uncertainties of our study, arriving at a final precision of 7\\% on the mean cluster mass.  In subsequent papers, we will incorporate the weak-lensing mass measurements into a self-consistent cosmological framework \\citep{mae10} in order to determine improved cosmological constraints on cosmology and key astrophysical scaling relations.  Accurate and precise absolute calibration of cluster masses will be a critical requirement for future studies aimed at constraining cosmological parameters with galaxy clusters. The methodology introduced in this series of papers should be straightforwardly adaptable to other projects, utilizing optical survey data and/or targeted follow-up observations of clusters."
    },
    "1208/1208.0304_arXiv.txt": {
        "abstract": "{It is unknown how far dust growth can proceed by coagulation. Obstacles to collisional growth are the fragmentation and bouncing barriers. However, in all previous simulations of the dust-size evolution, only the mean collision velocity has been considered, neglecting that a small but possibly important fraction of the collisions will occur at both much lower and higher velocities.} {We study the effect of the probability distribution of impact velocities on the collisional dust growth barriers.} {We assume a Maxwellian velocity distribution for colliding particles to determine the fraction of sticking, bouncing, and fragmentation, and implement this in a dust-size evolution code. We also calculate the probability of growing through the barriers and the growth timescale in these regimes.} {We find that the collisional growth barriers are not as sharp as previously thought. With the existence of low-velocity collisions, a small fraction of the particles manage to grow to masses orders of magnitude above the main population.} {A particle velocity distribution softens the fragmentation barrier and removes the bouncing barrier. It broadens the size distribution in a natural way, allowing the largest particles to become the first seeds that initiate sweep-up growth towards planetesimal sizes.} ",
        "introduction": "\\label{sec:introduction}  Primary accretion is the earliest stage of planet formation, where tiny, micrometer-sized dust grains in the protoplanetary disk grow into planetesimals of several kilometers in size. In the classic scenario, this happens by the way of incremental growth, in which sticking collisions lead to successively larger aggregates \\citep{1997Icar..127..290W, 2005A&A...434..971D}. As laboratory experiments \\citep{2008ARA&A..46...21B, 2010A&A...513A..56G} and numerical simulations \\citep{2010A&A...513A..57Z,2010A&A...513A..79B} have improved our understanding of the collision process and the dust evolution, it has become evident that incremental growth to form planetesimals cannot continue unhindered.  As particles grow, they decouple more and more from the surrounding gas, which increases their relative velocities. At the fragmentation barrier, collision energies are high enough to cause particle destruction, halting the dust growth at centimeter to meter sizes \\citep{2008A&A...480..859B}. \\cite{2010A&A...513A..57Z} also introduced the bouncing barrier, which stops the growth at even smaller sizes. In this case, the collision energies are too low to cause any particle destruction, but also too high for sticking, with growth-neutral bouncing events as the result.  \\cite{2012A&A...540A..73W} suggested a sweep-up scenario where the fragmentation barrier can be circumvented. They found that even though collisions between equal-sized particles generally lead to fragmentation, if the mass-ratio is large enough, growth of the larger particle can occur even at very high velocities. In this scenario, the growth initially stalls at the bouncing barrier, but if a small number of slightly larger 'seed' particles are introduced, they rapidly sweep up the smaller particles and grow to very large sizes. The growth barriers in this case limit the number of large particles and therefore reduce the number of destructive collisions among them. Exactly how the first seeds are formed is however still not clear.  All prior dust coagulation models have until now relied on the mean value to describe the velocity at which a collision occurs. In reality, the relative velocities between the particles that arises because of Brownian motion and turbulence does not take a single value, but is better represented by a probability distribution, owing to geometrical and stochastic effects. Here, we focus on turbulence since it is the dominating source of relative velocity between the small grains below the fragmentation barrier.  A general formula for the probability distribution function (PDF) of particle relative velocities is however unavailable, despite the efforts of many numerical and experimental works. There is tentative evidence that the PDF for particles with large Stokes numbers (${\\rm St} \\sim 1$, those that couple to the driving scales of the turbulence) is Maxwellian or close to it (\\citealt{2010MNRAS.405.2339C}; Dittrich et al., in prep.). However, at smaller sizes (where particles couple to the Kolmogorov scale) the PDF may be better characterized by wide, exponential tails (\\citealt{2000JFM...415..117W, 2010JFM...661...73P, Hubbard:2012ud}) Future numerical and analytical modelling is desired to refine and interpret these data. In this work, as a first step, we assume that turbulent velocities are Maxwellian distributed.  A velocity distribution allows some collisions to result in sticking where the average outcome would produce a bouncing or fragmentation event. This causes the barriers to blur out, and might allow for some lucky particles to just by sheer chance repeatedly experience only low-velocity collisions and grow to larger sizes than the main population.  In this letter, we show the effect of such a velocity distribution in a local dust-size evolution code, not only as a method for creating lucky larger particles, but also to see how it affects the general dust population. ",
        "conclusions": "\\label{sec:discussion}  We have found that the collisional growth barriers for dust grains are smoothed out and can even be overcome by virtue of a probability distribution of relative velocities among dust grains. Although improbable, sticky, low-velocity collisions can occur at sizes where the mean collisional velocity would lead to only bouncing or fragmentation.  To grow through the fragmentation barrier, a particle needs to be lucky and experience low-velocity collisions many times in a row, which causes a tail of larger particles to extend from the peak of the mass distribution. Assuming a Maxwellian velocity distribution, the luckiest particles can grow to around 50 times more massive than they would otherwise be.  The bouncing barrier is even more affected by the existence of a velocity distribution, and particles can grow to more than three orders of magnitude higher in mass, with the peak being shifted by two orders of magnitude. This occurs because low-velocity collisions lead to sticking, but even the higher velocities are low enough to only cause bouncing. This means that the growth can continue unimpededly until the average relative velocities have increased enough for the fragmentation barrier to start to become important. The bouncing barrier is therefore not a solid barrier at all, unless the growth timescale becomes too long because of the low sticking probability.  We have found that the low-velocity tail allows some lucky particles to grow beyond the bouncing and fragmentation barriers, to become the first seeds in the sweep-up scenario introduced by \\cite{2012A&A...540A..73W}. When the effect of fragmentation-with-mass-transfer is included, these seeds can sweep up the smaller particles trapped by the growth barriers, and start their growth towards planetesimal sizes."
    },
    "1208/1208.3900_arXiv.txt": {
        "abstract": "We present an analysis of selection biases in the \\msig\\ relation using Monte-Carlo simulations including the sphere of influence resolution selection bias and a selection bias in the velocity dispersion distribution. We find that the sphere of influence selection bias has a significant effect on the measured slope of the \\msig\\ relation, modeled as $beta_{intrinsic}=-4.69+2.22\\beta_{measured}$, where the measured slope is shallower than the model slope in the parameter range of $\\beta > 4$, with larger corrections for steeper model slopes. Therefore, when the sphere of influence is used as a criterion to exclude unreliable measurements, it also introduces a selection bias that needs to be modeled to restore the intrinsic slope of the relation. We find that the selection effect due to the velocity dispersion distribution of the sample, which might not follow the overall distribution of the population, is not important for slopes of $\\beta\\sim4$--6 of a logarithmically linear $M_{bh}-\\sigma$ relation, which could impact some studies that measure low (e.g., $\\beta<4$) slopes. Combining the selection biases in velocity dispersions and the sphere of influence cut, we find the uncertainty of the slope is larger than the value without modeling these effects, and estimate an intrinsic slope of $\\beta=5.28_{-0.55}^{+0.84}$. ",
        "introduction": "The relation between super-massive black hole (SMBH) masses, \\m, and their host galaxy velocity dispersion, \\vd, is one of the central themes in galaxy and AGN studies. The number of reliable SMBH estimates has more than quadrupled since the early stages of investigation \\citep[e.g.,][]{fm00,g00}, with the latest samples containing $\\sim64$ galaxies \\citep[e.g.,][]{graham11,mc11}. Even with an increase in sample size, the relation remains tight, e.g., with a scatter of only 0.43 dex \\citep{graham11}, tantalizingly suggestive of a strong link between SMBH and galaxy formation, and indicative of AGN feedback processes. Many theoretical predictions and models are based on this relation, describing the interaction between SMBH and the bulge of the galaxy through energy-driven winds \\citep[e.g.,][]{silk98} or momentum-driven winds \\citep[e.g.,][]{fabian99,k03}. The \\msig\\ relation has been used to support many theories connecting the central black hole to other host properties of the galaxy \\citep[e.g.,][]{ff05,adams03,mh08,bw06,dm08,onb08}, where connections as far as with dark matter halos have also been suggested \\citep[e.g.,][]{bs10}. There are also studies \\citep[e.g.,][]{gultekin09,graham11,beifiori12} that examine the possibility of the \\msig\\ relation varying with the morphology of the galaxy and evolution with redshift \\citep[e.g.,][]{p06}.  The slope and normalization of the \\msig\\ relation are important parameters for constraining the nature of the feedback mechanism from the black hole. Various models and simulations generally predict $M_{bh} \\propto \\sigma^5$ \\citep[e.g.,][]{silk98} if the feedback is in the mechanical form and $M_{bh} \\propto \\sigma^4$ if the feedback is dominated by momentum exchanges \\citep[e.g.,][]{fabian99,k03,g04,m05}. The normalization of the relation constrains the feedback efficiency, where $\\eta \\sim 0.05$ is generally required for energy feedback models and $\\eta \\sim 1$ for momentum feedback models. Thus, momentum feedback models require that the majority of the black hole growth occurs in the obscured phase, which can be tested by future hard X-ray surveys. Energy feedback models suffer from cooling problems where the feedback energy is lost by Compton scattering with photons from AGN or other radiative cooling mechanisms \\citep[e.g.,][]{k03,sn10}. Even the 5\\% feedback efficiency for energy feedback models is difficult to achieve based on current observations of AGN winds. Observationally, there are a number of attempts to constrain the \\msig\\ relation (Table~1), where the slope is measured between 4--5 but with large enough errors and variety in measurements that it is unclear which feedback model drives the relation. The average slope of major studies {is $\\beta=4.7$, which is inclined towards an energy-driven feedback scheme, but the samples used to measure the $M_{bh}-\\sigma$ relation are still under close scrutiny.  Considering the importance of the \\msig\\ relation, it is essential to thoroughly investigate the relation. Authors have previously examined the effect of the uncertainties in black hole and velocity dispersion measurements \\citep[e.g.,][]{nfd06,mf01}, as well as explored selection biases in choosing samples of quasars versus nearby AGN \\citep[][]{sw11}. The treatment of errors arguably can change the measured slope of the \\msig\\ relation, and as \\citet{mf01} discuss, the limited sample size can also impact the relation. Recently, questions were raised whether the fitting results of the \\msig\\ relation are affected by selection effects and even whether the relation itself could be an artifact of selection effects. Studies have tended to use resolution of the sphere of influence, \\ri, around the SMBH as criteria for inclusion of a galaxy in their sample as a proxy for reliable mass measurements \\citep[e.g.,][]{ff05,h08}, while \\citet{gultekin09} argue that no such criteria is needed. Restricting samples to galaxies where the sphere of influence, \\ri$=G$\\m$/\\sigma^2$ \\citep{p72}, is resolved provides a lower bound to what is observable in the \\msig\\ plane. Therefore, directly fitting these filtered samples can result in parameters different from those of an intrinsic relation \\citep[e.g.,][]{gultekin09}. Furthermore, a relation measured from a sample using this criteria might even arise simply from setting the lower limit for \\msig\\ data pairs. \\citet{b10} addresses this issue by applying both the \\ri\\ resolution cut and the observed upper limit plus scatter in the \\msig\\ plane and can reproduce the \\msig\\ relation in simulations. In particular, the author uses observed \\vd\\ and simulated \\m, evenly distributed logarithmically over $10^1$\\msun\\ to $10^{10}$\\msun, and makes an \\ri\\ cut based on an Hubble Space Telescope (HST) resolution of 0\\sarc1 and uses the upper scatter of the \\msig\\ relation itself to create an upper bound. This simple approach provides a physically different interpretation of the observed \\msig\\ relation: instead of a correlation, there is an envelope that sets an upper bound for \\m\\ given \\vd, which may represent a more realistic model for black hole growth \\citep{k10}. The lower bound therefore only arises from selection effects in this model. However, this model predicts a distribution of detection rates of black hole masses that are arguably not supported bya observations \\citep{gultekin11}.  Aside from the sphere of influence selection bias, there can be another selection effect in the \\vd\\ distribution that is also worth exploring. Ideally, the relation between \\m\\ and \\vd\\ is best constrained if the individual parameters' measured distributions in a small sample follow the general distributions from a large sample. Although it is impossible to compare to the true \\m\\ distribution, we find differences between the observed \\vd\\ distribution in the \\msig\\ sample and the distribution from a large sample of galaxies. As shown in our simulation, this difference can also change the best-fit values for model parameters, especially when the underlying relation is non-linear or skewed. In this paper, we model both the effects of the sphere of influence resolution criteria and selection effects caused by \\vd\\ distributions. In addition, continuing the argument of \\citet{b10}, we extend the question of whether an upper bound can be naturally produced for the observed \\msig\\ distribution by simulating two more general mass distributions of SMBH masses independently of the \\vd\\ of galaxy bulges. Finally, most previous studies use various direct linear regression techniques to measure the parameters of the \\msig\\ relation, which make it difficult to model various selection effects, and can certainly impact the final measured slope \\citep{mf01}.  In this paper, we use a Bayesian Monte Carlo approach to robustly measure the parameters of the \\msig\\ relation including selection effects.  In Section 2, the samples used are described. Section 3 discusses methods and results, followed by conclusions in Section 4. Throughout this paper, unless otherwise noted, we assume a cosmology of $H_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$=0.27, and $\\Omega_{\\Lambda}$=0.73. ",
        "conclusions": "Using a Bayesian Monte-Carlo analysis, we constrain the parameters of the \\msig\\ relation as $\\alpha = 8.07_{-0.10}^{+0.08}$, $\\beta = 5.28_{-0.55}^{+0.84}$, and $\\epsilon = 0.45_{-0.07}^{+0.08}$. The main results of this paper are: (1) the directly measured slope of an \\msig\\ relation is shallower than the intrinsic slope because of selection effects in \\ri\\ and \\vd , and (2) the uncertainties on the slope of the relation are likely underestimated in previous measurements.  We additionally conclude that the \\msig\\ relation cannot be successfully reproduced with either power law or exponential cutoff mass distributions coupled with \\ri\\ and \\vd\\ selection effects, and that it is statistically much more likely that the relation is intrinsic rather than observed.  Comparing the two selection effects shows clearly that the sphere of influence selection effect dominates, which will decrease the slope if the intrinsic slope is steep, e.g., $\\beta \\sim 10$, or increase the slope if the intrinsic slope is shallow, e.g., $\\beta \\sim 1$, with the crossing point at $\\beta \\sim 4$ (Figure~\\ref{fig:inout}). The \\ri\\ selection resembles a parabolic curve in the logarithmic \\msig\\ plane, which can be approximated with a linear cut of $\\beta \\sim 3.5$. For a steep intrinsic \\msig\\ relation, more objects with lower masses will be cut if one applies the \\ri\\ resolution selection criteria, resulting a shallower relation. The situation will reverse for very shallow intrinsic slopes (e.g., $\\beta \\sim 1$), where the simulated slope will be steeper than the intrinsic input slope after applying selection effects. However, since the likelihood function for these cases approaches zero, this effect is not important. For the power law and exponential cutoff mass distributions, our models assume more galaxies with smaller \\m\\ given the range of $\\Gamma$ used in the simulations.  However, in the high \\vd\\ range, the \\ri\\ selection will trim most of the low mass objects, and in the low \\vd\\ range, only a few objects have measurements in \\m\\ and it is more likely for the mass to be small based on our assumptions. Even though we are able to rule out these two scenarios, it is still clear that selection effects can significantly increase the chance that these intrinsic power law models would be consistent with the observed sample.  The \\vd\\ selection effect can be important if the intrinsic relation between \\m\\ and \\vd\\ is non-linear as in our power law models.  For linear models, the \\vd\\ selection effect can also be important if the intrinsic slope is extremely deep, $\\beta \\sim 10$, or shallow, $\\beta \\sim 1$, where the effect can modify the slope by $\\Delta \\beta \\sim 1$.  Fortunately, within the range of intrinsic slopes of $\\beta = 4$--6, the \\vd\\ selection effect is not important.  The argument for using an \\ri\\ selection criteria is that black hole mass estimates can be unreliable \\citep[e.g.,][]{ff05}. Ensuring that the sphere of influence is resolved is thought to more accurately measure the \\msig\\ relation. In this paper, we find that this will cause selection effects that impact estimating the best-fit parameters, especially the slope $\\beta$.  Based on our simulations, we find \\begin{equation} \\beta_{intr} = -4.69 + 2.22\\times\\beta_{measured}, \\end{equation} where $\\beta_{intr}$ is the intrinsic slope and $\\beta_{measured}$ is the directly measured slope without modeling the selection effects. In the relevant parameter space $\\beta = 4$--6, the \\ri\\ selection effect decreases the measured slope. This is consistent with some previous results \\citep[e.g.,][]{fm00}. The authors measured the \\msig\\ relation twice, once with secure measurements and the sphere of influence resolved, and once for their entire sample. They found that the slope decreases from $5.81\\pm0.43$ to $4.80\\pm0.54$ when including only galaxies where \\ri\\ was resolved. \\citet{gultekin09} used Monte-Carlo simulations to show that an \\ri\\ selection criteria can bias measurements of the zero point, slope, and scatter. They consider three different scenarios using synthetic data sets. In the first scenario, they require that \\ri\\ be resolved based on the values of $M_{bh}$ and \\vd\\ in the data set. For the next scenario, they use \\vd\\ to calculate the expected $M_{bh}$, find the radius of the sphere of influence using \\vd\\ and the expected mass, and then cut any data pairs where \\ri\\ is not resolved. The last scenario uses an observed sample of \\vd\\ with simulated masses. The results here agree with their first scenario, as the slope decreases with selection effects. However, we find an opposite bias in the slope for the third scenario, which is closest to the simulations here as it uses observed \\vd. The differences might stem from the fact that the \\citet{gultekin09} scenario uses observed distances and instrumental resolution, which give different values for \\ri, rather than assuming a standard resolution for all data points, as done here. In addition, \\citet{gultekin09} studies a sample yielding a low slope, which, based on our simulation, has a smaller correction on $\\beta$ (Figure~\\ref{fig:inout} and Equation~2). To test the robustness of our conclusions, we run our Bayesian Monte-Carlo code with other recent samples \\citep{gultekin09,mc11, beifiori12}. We list our corrections on the slope, $\\Delta \\beta$, in Table~\\ref{slopes}. We find decreases of the measured slope after applying selection effects in all the samples, however, with larger corrections for larger measured slopes.  This is consistent with the results of our simulation.  There are a number of measurements of the \\msig\\ relation, but we do not find an analysis that fully models various selection effects. The fitting results among  different groups are not consistent, due to the mass measurements, sample selections, or linear regression fitting techniques (about $\\Delta \\beta \\sim 0.2$). Using the sample of \\citet{graham11} but filtering out the two galaxies not satisfying our simulation resolution of 0\\sarc08, we find a best estimated slope of $\\beta = 5.28_{-0.55}^{+0.84}$ by modeling the selection effects. Although this is in agreement with studies that measure higher values for the slope (e.g., \\citet{ff05}, $\\beta=5.81\\pm0.43$; \\citet{graham11}, $\\beta=5.13\\pm0.34$; \\citet{gr08}, $\\beta=5.52\\pm0.40$), this consistency needs to be further evaluated. For example, our analysis indicates that directly fitting the \\citet{graham11} sample should yield a slope of $\\beta \\sim 4.5$ (Table~2), and indeed, we find a slope of $\\beta \\sim 4.6$ by directly performing a linear regression on the volume-limited \\citet{graham11} sample with which we compare our simulations. Applying the \\ri\\ selection is important to keep the mass measurement reliable; however, we also need to model this effect to restore the intrinsic slope of the relation. For groups measuring a shallower slope, $\\beta \\sim 4$, the selection effect is small based on our simulation, and for measurements of steeper slopes $\\beta \\sim 5$, the selection effect is larger. Therefore, it appears the difference in the measurements of slopes ($\\sim 4$ vs. $\\sim 5$) among different groups arises mainly from the basic mass measurements.  An \\msig\\ relation has far reaching implications, as the slope of the relation can illuminate important physical processes connecting a central black hole to the host galaxy. The results presented here suggest that the slope is likely higher than is measured due to \\ri\\ and \\vd\\ selection effects, and closer to the energy-driven winds scheme as predicted by \\citet[$M\\propto\\sigma^5$;][]{silk98} rather than the momentum-driven winds predicted by \\citet[$M\\propto\\sigma^4$;][]{fabian99}. This feedback mechanism can be realized by quasar winds, especially in broad absorption line quasars (BALQSOs). The efficiency depends linearly on the covering fraction of BALQSOs, and recent studies suggest that the intrinsic fraction of BALQSOs in quasars is at least two times higher than those fractions measured in optical surveys \\citep[e.g.,][]{dai08,s08,allen11}.  For low-ionization BALQSOs (LoBALs),  including ones exhibiting iron absorption (FeLoBALs), the intrinsic fractions are estimated to be 5--7 times larger than the values obtained from optical surveys \\citep{dai10b}.  Although the intrinsic fractions of LoBALs are still small, the column density in these wind can be orders of magnitude higher than other quasar winds.  Using more robust measurements of the column densities of FeLoBALs, it is estimated that FeLoBALs can provide kinetic feedback efficiency $\\dot{E}_{kin}/L_{bol}$ of a few percent \\citep[e.g.,][]{m11}, consistent with the requirement for the feedback efficiency of the \\msig\\ relation \\citep[e.g.,][]{silk98}. Although the energy feedback model can suffer from radiation loss \\citep[e.g.,][]{k03,sn10}, it is possible that the feedback process is a multi-phase process, where the AGN feedback is just the first phase \\citep[e.g.,][]{he10,sn10}.  Subsequent feedback from supernovae and stellar winds will provide additional energy.  The slope of the predicted \\msig\\ relation is still $\\propto \\sigma^5$ if the feedback is dominated by energy exchanges \\citep{sn10}.  An alternative model that is not studied here is the combination of an upper envelope and selection effects causing us to see only the upper portion of the \\msig\\ plane \\citep{b10}, as physically interpreted by \\citet{king10}. However, \\citet{gultekin11} argue that these upper-limit models predict detection rates of black hole masses that are not observed. The lower portion of the \\msig\\ plane will be probed as technology and techniques advance enough to be able to resolve low mass black holes. Continuing to expand the observed samples with reliable SMBH mass estimates is essential, but until we reach the point where selection effects are negligible, they must be well modeled when analyzing a sample.  Here we have only tested a limited number of hypotheses against the current data, and therefore our conclusions are limited to those hypotheses. There are other factors that are not explored: for example, many studies differentiate galaxies by morphology when measuring the \\msig\\ relation (e.g., Graham et al. 2011; Hu 2008; Beifiori et al. 2012). While there are tantalizing hints that the \\msig\\ relation may be different for different galaxy populations, the sample sizes are still too limited to be conclusive. This paper provides a formalism that is capable of testing a large number of hypotheses combined with selection effects, a useful springboard for future, deeper studies."
    },
    "1208/1208.4711_arXiv.txt": {
        "abstract": "{A systematic study of RR Lyrae stars is performed based on a selected sample of 655 objects in the Large Magellanic Cloud with observation of long span and numerous measurements by the Optical Gravitational Lensing Experiment III project.  The Phase Dispersion Method and linear superposition of the harmonic oscillations are used to derive the pulsation frequency and variation properties. It is found that there exists an Oo I and Oo II dichotomy in the LMC RR Lyrae stars. Due to our strict criteria to identify a frequency, a lower limit of the incidence rate of Blazhko modulation in LMC is estimated in various subclasses of RR Lyrae stars.  For fundamental-mode RR Lyrae stars,  the rate 7.5$\\%$ is smaller than previous result. In the case of the first-overtone RR Lyr variables, the rate 9.1$\\%$ is relatively high. In addition to the Blazhko variables, fifteen objects are identified to pulsate in the fundamental/first-overtone double mode. Furthermore, four objects show a period ratio around 0.6 which makes them very likely the rare pulsators in the fundamental/second-overtone double-mode.} ",
        "introduction": "RR Lyrae stars (RRLS) are pulsating variables on the horizontal branch in the H-R diagram. They have short periods of 0.2 to 1 day and low metal abundances Z of 0.00001 to 0.01. Usually they can be easily identified by their light curves and color-color diagrams \\citep{2011RAA....11..833L}. RRLS is famous for the \"Blazhko effect\" \\citep{1907AN....175..325B}, a periodic modulation of the amplitude and phase in the light curves, which is still a mystery in theory today. The main photometric feature of the Blazhko effect is that the frequency spectra of the light curves are usually strongly dominated by a symmetric pattern around the main pulsation frequency $f_{\\rm 0}$, i.e., $k f_{\\rm 0}$ and $k f_{\\rm 0}\\pm f_{\\rm BL}$ where $f_{\\rm BL}$ is the modulation frequency and $k$ is the harmonic number \\citep{2009AIPC.1170..261K}. Moreover, the higher-order multiplets such as quintuplets and higher, i.e. multiplets $k f_{\\rm 0}\\pm l f_{\\rm BL}$, are now also attributed to the Blazhko effect \\citep{2011arXiv1106.4914B}. On the other hand, the amplitudes of modulation components are usually asymmetric so that one side could be under the detection limit in highly asymmetric cases, which may lead to the asymmetric appearance of the frequency spectra. Several models are proposed to explain this effect, including the non-radial resonant rotator/pulsator \\citep{1994A&A...291..481G}, the magnetic oblique rotator/pulsator  \\citep{2000ASPC..203..299S}, 2:1 resonance model \\citep{1980AcA....30..393B},  resonance between radial and non-radial mode \\citep{2004AcA....54..363D}, 9:2 resonance model \\citep{2011ApJ...731...24B} and convective cycles model \\citep{2006ApJ...652..643S}. However, none of them is able to interpret all the observational phenomena about the Blazhko effect.  A systematic study of RRLS helps to understand their nature, such as the incidence rate of various pulsation modes, the distribution of modulation frequency and amplitude and the dependence of the variation properties on environment. This has been performed on the basis of some datasets with large amount of data for variables. The early micro-lensing projects, MACHO \\citep{2003ApJ...598..597A} and OGLE \\citep{2008AcA....58..163S} that surveyed mainly LMC, SMC and the Galactic bulge, and the variables-oriented all-sky survey project, ASAS \\citep{2007MNRAS.377.1263S},  have given particularly great support to such study. Using the MACHO data, \\cite{2000ApJ...542..257A} and \\cite{2006A&A...454..257N} analyzed the frequency of 1300 first overtone RRLS in LMC with an incidence rate of the Blazhko variables of 7.5\\%. Meanwhile, \\cite{2003ApJ...598..597A} made a frequency analysis of 6391 fundamental mode RRLS in LMC that resulted in an incidence rate of the Blazhko variables of 11.9\\%.  With the OGLE-I data, \\cite{2003A&A...398..213M} searched for multiperiodic pulsators among  38 RRLS in the Galactic Bulge. \\cite{2003AcA....53..307M} made a complete search for multi-period RRLS from the OGLE-II database, while \\cite{2006ApJ...651..197C} presented a catalogue of 1888 fundamental mode RRLS  in the Galactic bulge from the same database.  Recently, \\cite{2008CoAst.157..345M} conducted a systematic search for multiperiod RRLS in $\\omega$ Centauri, a globular cluster, and found the incidence rate of Blazhko modulation pretty high, about 24\\% and 38\\% for the fundamental and first-overtone RRLS respectively. \\cite{2009MNRAS.400.1006J} got a 47\\% incidence rate in the dedicated Konkoly survey sample of 30 fundamental-mode RRLS in the Galactic field.  \\cite{2010ApJ...713L.198K} also claimed at least 40\\% RRLS in the  Kepler space mission sample of 28 objects exhibiting the modulation phenomenon.  The OGLE-III database released in 2009 \\citep{2009AcA....59....1S} contained 24906 light curves that are preliminarily classified as RRLS in the Large Magellanic Cloud, the ever largest sample of RRLS. This database covers a time span of about 10 years that makes a large-scale analysis of the RRLS variation possible. \\cite{2009AcA....59....1S} analyzed the basic statistical features of RRLS in the LMC and divided them into 4 subtypes: RRab, RRc, RRd and RRe. With this sample of RRLS, the study of the structure of LMC was developed. \\cite{2009ApJ...704.1730P} investigated the structure of the LMC stellar halo; \\cite{2009A&A...503L...9S} found that RRLS in the inner LMC trace the disk and probably the inner halo; \\cite{2010MNRAS.408L..76F} established a small but significant radial gradient in the mean periods of Large Magellanic Cloud (LMC) RR Lyrae variables. The data of RRLS in SMC and the Galactic bulge are also  released \\citep{2010AcA....60..165S,2011AcA....61....1S} and some work is also done with those data \\citep{2011arXiv1107.3152P}. It is worth to note that these work mainly deals with the structure of LMC other than the RR Lyr variables.  In this paper, we focus the study on the RR Lyr variables themselves based on the released OGLE-III database. With the PDM and Fourier fitting methods, we make a precise systematic frequency analysis of carefully selected 655 RRLS  in LMC, and a detailed classification of the RRLS in the LMC based on which to discuss the incidence rate of the Blazhko modulation in various pulsation modes. The data and the sample are illustrated in section 2,  the method is introduced in section 3,  the detailed classification of RRLS in LMC in section 4 and  discussion in section 5. ",
        "conclusions": "The incidence rates of each subclass are shown in Table~\\ref{tab9}. The majority, 85\\%, is single-mode pulsators. The RRLS exhibiting the Blazhko effect (sum of RR-BL1 and RR-BL2) is the second most numerous group. With 52 RR-BL stars, they consist 7.9\\% of the sample. The incidence rates of Blazhko stars are compared with previous results in Table~\\ref{tab10}. As our identification of a frequency is quite strict, the percentages in our sample should be taken as the lower limit of the Blazhko incidence rate.   For RRLS in LMC, the Blazhko variables (sum of RR-BL2 and RR-BL1 stars) occur less frequently in RR0 (7.5\\%) than in RR1 stars (9.1\\%). For RR1 stars, we can see an increasing trend of the incidence rates from 2.0\\% and 7.5\\% in previous work to 9.1\\% in present work, this can be explained by the longer time span and more precise data. But this trend does not show in RR0 stars although the data has been improved as for RR1 stars. To further analyze the reason, it is found that the incidence rate of RR-BL1 is comparable to previous work, the rate is 6.7\\% for RR0-BL1 and 5.1\\% for RR1-BL1, compared to 6.5\\% \\citep{2003ApJ...598..597A} and 3.5\\% \\citep{2006A&A...454..257N}. But in regarding to the RR-BL2 stars, the incidence rate of RR1-BL2 is comparable to the work before, with 5.1\\% to 3.5\\%. Meanwhile the incidence rate of RR0-BL2 is abnormally low, with only 0.8\\% compared to 5.4\\% in \\cite{2003ApJ...598..597A}. As mentioned in previous section, the reason may lie on our very strict criteria to accept a frequency which can move a BL2 star into a BL1 star in the case of high asymmetry of the amplitude in the two side frequencies. This explanation finds support in the fact that the RR0-BL2 stars have the amplitude ratio $A_{+}/A_{-}$ not far from unity.  Another possible reason is that, among RR0-BL1 stars, 75\\% (24 out of 32) have $f_+$, and all the four RR0-BL2 stars have $A_+ >A_-$.  In the work of \\cite{2002ASPC..259..396K}, 80\\% of RR0-BL1 stars were found to have $A_+$ component. We suspect that for RR0-BL stars, there maybe some unknown effect to make $A_+$ much larger than $A_-$, which made a lot of $A_-$ components missing. This effect does not appear in RR1 stars, for example, \\cite{2000ApJ...542..257A} found  37\\% RR1-BL1 stars have $A_-$ components, and in our sample only 43\\% RR1-BL2 stars have $A_+ > A_-$. Based on Eq.(45) from \\cite{2011arXiv1106.4914B}, most of the RR0 stars have $\\pi < \\phi _m < 2 \\pi$; while for those RR1 stars, $\\phi _m$ is evenly distributed between zero and $2 \\pi$.  People used to believe that the incident rates of the Blazhko variables are lower in LMC than in the Galaxy. But this only suits for RR0 stars, may not be true for RR1, as the work of \\cite{2006A&A...454..257N} already suggested. From our work, with the long time span of observation, the Blazhko incidence rate for RR1 stars is larger in LMC than in the Galaxy bulge, as the rates are 5.1\\% and 4.0\\% for RR1-BL1 and RR1-BL2 respectively in our LMC sample in comparison with the 3.1\\% and 1.5\\% from \\cite{2003A&A...398..213M} or 2.9\\% and 3.9\\% from \\cite{2003AcA....53..307M} for the bulge RRLS sample. On the other hand, both the improvement of the observational precision and the extension of observational span increase the possibility to detect the Blazhko effect. Based on the data from the \\emph{Kepler} mission \\cite{2010ApJ...713L.198K} and  the Konkoly Blazhko  Survey  \\citep{2009MNRAS.400.1006J}, the incidence rate can exceed forty percent, but they are small samples with no more than 30 objects. Thus such comparison may not be conclusive as the observation and the analysis techniques are not uniform, and the samples are very different. According to these observations of LMC, SMC, bulge and $\\omega$ Cen, there is no clear relation between the incidence rate and metallicity. What causes the difference in the Blazhko incidence rate in different environments is unclear, which could be part of the difficulty in understanding the mechanism for the Blazhko effect.  \\begin{table} \\caption{Statistical results of the RRLS classification.} \\centering \\begin{tabular}{ccrrcr} \\hline\\hline                        %  type\t&\tShort description\t&\tnumber &\tpercent\t&\tsubtype \t&\tnumber\t\\\\ \\hline RR-SG & Single-period & 556 & 84.9\\%\t&\tRR0-SG\t&\t424\t\\\\ &\t\t&\t\t&\t\t&\tRR1-SG\t& 132\t\\\\ \\hline RR01\t&\tFU/FO double mode \t&\t15\t&\t2.3\\%\t& &\t\t\\\\ \\hline RR02\t&\tFU/SO double mode \t&\t4\t&\t0.6\\%\t& &\t\t\\\\ \\hline {RR-BL}\t&\t{One close component}\t& {52}\t&\t{7.9\\%}\t&\tRR0-BL1\t& 32\t\\\\ &\t\t&\t\t&\t\t&\tRR1-BL1\t& 9\t\\\\ &\t{Two symmetric components}\t& &\t\t&\tRR0-BL2\t &\t4\t\\\\ &\t\t&\t\t&\t\t&\tRR1-BL2\t& 7\t\\\\ \\hline {RR-MC}\t&\t{Multiple close components} &\t {4}\t&\t{0.6\\%}\t&\tRR0-MC &\t1\t\\\\ &\t\t&\t\t&\t\t&\tRR1-MC\t& 3\t\\\\ \\hline {RR-PC}\t&\t{Period changing}\t& {20}\t&\t{3.1\\%}\t&\tRR0-PC\t& 11\t\\\\ &\t\t&\t\t&\t\t&\tRR1-PC\t& 9\t\\\\ \\hline RR-D1\t&\tSecond frequency at unity \t&\t2\t&\t0.3\\% &\tRR0-D1\t&\t2\t \\\\ \\hline RR-?\t&\tMysterious double mode \t&\t2\t&\t0.3\\%\t& RR0-others\t&\t2\t \\\\ \\hline \\end{tabular} \\label{tab9} \\end{table}   \\begin{table} \\centering \\caption{The incidence rates of Blazhko effect in the sample compared to previous results. The references in the table are: (a): \\cite{2003A&A...398..213M}, (b): \\cite{2003ApJ...598..597A}, (c): \\cite{2003AcA....53..307M}, (d): \\cite{2006ApJ...651..197C}, (e): \\cite{2008CoAst.157..345M}, (f): \\cite{2000ApJ...542..257A}, (g): \\cite{2006A&A...454..257N}, \\textbf{(h): this work}.} \\begin{tabular}{crrrrrr} \\hline\\hline Ref.\t&\tBulge(a)\t&\tBulge(c)\t& Bulge(d)\t&\t LMC(b)\t&\t $\\omega$ Cen (e)\t & \\textbf{LMC(h)}\t\\\\ \\hline RR0\t&\t150\t&\t1942\t&\t1888\t&\t6391\t& 70\t&\t\t\\textbf{478}\t \\\\ RR0-BL1\t&\t16.7\\%\t&\t12.5\\%\t&\t8.8\\%\t&\t6.5\\%\t& 4.3\\%\t&\t\t \\textbf{6.7\\%}\t \\\\ RR0-BL2\t&\t6.0\\%\t&\t7.4\\%\t&\t14.9\\%\t&\t5.4\\%\t& 18.6\\%\t&\t\t \\textbf{0.8\\%}\t \\\\ RR0-MC\t&\t\t&\t4.4\\%\t&\t3.5\\%\t&\t0.3\\%\t& &\t\\textbf{0.2\\%}\t \\\\ RR0-PC\t&\t\t&\t\t&\t6.3\\%\t&\t2.9\\%\t& &\t\\textbf{2.3\\%}\t\\\\ \\hline  \\hline Ref.\t&\tBulge(a)\t&\tBulge(c)\t&\tLMC(f)\t& LMC(g)\t&\t$\\omega$ Cen (e)\t &\t\t\\textbf{LMC(h)}\t\\\\ \\hline RR1\t&\t65\t&\t771\t&\t1327\t&\t1332\t& 81\t&\t\t\\textbf{177}\t\\\\ RR1-BL1\t&\t3.1\\%\t&\t2.9\\%\t&\t1.8\\%\t&\t3.5\\%\t& 21.0\\%\t& \\textbf{5.1\\%}\t \\\\ RR1-BL2\t&\t1.5\\%\t&\t3.9\\%\t&\t0.2\\%\t&\t4.0\\%\t& 4.9\\%\t& \\textbf{4.0\\%} \\\\ RR1-MC\t&\t\t&\t5.3\\%\t&\t0.4\\%\t&\t1.0\\%\t& &\t\t \\textbf{1.7\\%}\t\\\\ RR1-PC\t&\t\t&\t\t&\t10.6\\%\t&\t14.0\\%\t& &\t\t\\textbf{4.5\\%}\t \\\\ \\hline \\end{tabular} \\label{tab10} \\end{table}"
    },
    "1208/1208.3846_arXiv.txt": {
        "abstract": "Most stars in galactic disks are believed to be born as a member of star clusters or associations. Star clusters formed in disks are disrupted due to the tidal stripping and the evolution of star clusters themselves, and as a results new stars are supplied to the galactic disks. We performed $N$-body simulations of star clusters in galactic disks, in which both star clusters and galactic disks are modeled as $N$-body (``live'') systems, and as a consequence the disks form transient and recurrent spiral arms. In such non-steady spiral arms, star clusters migrate radially due to the interaction with spiral arms. We found that the migration timescale is a few hundreds Myr and that the angular momentum changes of star clusters are at most $\\sim 50$\\% in 1 Gyr. Radial migration of star clusters to the inner region of galaxies results in a fast disruption of the star clusters because of a stronger tidal field in the inner region of the galaxy. This effect is not negligible for the disruption timescale of star clusters in galactic disks. Stars stripped from clusters form tidal tails which spread over 1--2 kpc. While the spatial distribution of tidal tails change in a complicated way due to the non-steady spiral arms, the velocity distribution conserve well even if the tidal tails are located at a few kpc from their parent clusters. Tidal tails of clusters in galactic disks might be detected using velocity plots. ",
        "introduction": "Star clusters are one of the fundamental building blocks of galactic disks because most stars are formed in star clusters \\citep{2003ARA&A..41...57L}. In disk galaxies, new clusters born in galactic disks travel in their host disks experiencing disruptions and supply new stars to the disks. The disruption of clusters is caused by the tidal force from their host galaxy and also the internal evolution of star clusters themselves such as dynamical evolution, mass loss due to the stellar evolution, and gas expulsion \\citep{2000MNRAS.318..753F, 2003MNRAS.340..227B,2004AJ....127.2753D,2005AJ....129.1906C, 2005A&A...441..117L,2008MNRAS.384.1231B,2011MNRAS.413.2509G}.  Non-axisymmetric structures in galactic disks, such as spiral arms \\citep{2007MNRAS.376..809G} and bars \\citep{2012MNRAS.419.3244B}, have also been expected to play important roles for the dynamical disruption of clusters. For example, \\citet{2007MNRAS.376..809G} investigated the effect of spiral-arm passages on the evolution of star clusters assumed to rotate in a fixed pattern speed, as in the stationary density wave theory \\citep{1964ApJ...140..646L,1996ssgd.book.....B}. However, self-consistent simulations of galactic disks have shown that self-excited spiral arms are not stationary regardless of the existence of gas (or some kind of dissipation) \\citep{1984ApJ...282...61S,2002MNRAS.336..785S,2003MNRAS.344..358B, 2009ApJ...706..471B,2011ApJ...730..109F,2011MNRAS.410.1637S, 2011ApJ...735....1W,2012MNRAS.421.1529G,Baba+2012}. Such non-steady spiral arms do not have a single pattern speed but roughly follow the galactic rotation \\citep{2011ApJ...735....1W,2012MNRAS.421.1529G}. Therefore, these arms scatter stars everywhere in the disk by the co-rotation resonance \\citep{2002MNRAS.336..785S,2012MNRAS.421.1529G,Baba+2012}. Non-steady spiral arms change the gravitational fields around star clusters chaotically rather than periodically as in the stationary density waves.   In order to know the dynamical evolution of star clusters in galactic disks with non-steady spiral arms, we need to model both star clusters and galactic disks as $N$-body (``live'') systems. Such self-consistent $N$-body simulations are technically more difficult than those with rigid potential disks because the dynamical timescale of star clusters is much shorter than that of galactic disks and a large number of particles are required for the modeling of the disks \\citep{2011ApJ...730..109F}. We solve this problem using a direct-tree hybrid code, Bridge \\citep{2007PASJ...59.1095F}.   In this letter, we perform self-consistent $N$-body simulations of star clusters in live disks using Bridge and demonstrate that the angular-momentum exchange between star clusters and spiral arms causes the radial migration of star clusters of a few kpc from their initial galacto-centric radii. The migration timescale is shorter than the galactic rotation timescale, i.e., a few hundred Myr. The radial migration causes the tidal disruption of star clusters bringing them to closer to the galactic center. Star clusters lose their mass in their perigalacticon passage, and their tidal tails spread over a few kpc. We also find that tidal-tail stars stay close to their parent clusters in their velocity space even if they are already a few kpc from the parent cluster. The tidal tail of young clusters might be detectable using their velocities. ",
        "conclusions": "We performed a series of $N$-body simulations of star clusters in live stellar disks with multiple spiral arms. In these simulations, both galactic disks and star clusters are modeled as $N$-body systems and integrated self-consistently. Our results show that star clusters migrate radially a few kpc in the time scale of their orbital period in the disk (a few 100 Myr) because of the angular-momentum exchange with transient spiral arms. The angular-momentum change of the clusters is at most $\\sim 50$\\% of the initial angular momentum within 1 Gyr.  In the case of non-steady transient spiral arms, the radial migration strongly affects the tidal disruption of star clusters because star clusters lose more mass when they approach the galactic center due to the smaller Jacobi radii. This kind of disruption mechanism does not appear in stationary density waves. The heating due to the non-steady (corotating) spiral-arm passage would not be as strong as that by the density-wave spiral arms, because the adiabatic change of the energy due to the slow passages of spiral arms suppresses the heating per passage, and  the number of spiral passages is quite few (see the corotation case in Figure 8 in \\citet{2007MNRAS.376..809G}). Furthermore, the radial migration of star clusters can carry stars far from their original orbital radii and finally the distribution of stars would be much wider than those in the case of the density waves \\citep{2010ApJ...713..166B}.   With transient spiral arms, star clusters and their tidal tails tend to stay in or close to spiral arms. The shape of tidal tails of clusters change in a complicated way in time compared to those in a smooth tidal field like a halo potential. When a star cluster approach the galactic center, star clusters move along a spiral arm as is the case of stars and gas moving in spiral arms. In this phase, the tidal tails also spread along the spiral arm. During the apogalacticon passage, on the other hand, the tidal tails are compressed and distribute around 1 kpc from the cluster even though they are unbound. The tidal tails of clusters might be detectable even if they spread over a few kpc, because the tidal-tail stars still remain very close to the cluster in velocity space after they become unbound. If we know the velocity of stars, we might be able to detect the tidal tails of star clusters in the Galactic disk using a future astrometry such as {\\it Gaia} and {\\it JASMINE}."
    },
    "1208/1208.1072_arXiv.txt": {
        "abstract": "\\baselineskip=0.6 cm \\begin{center} {\\bf Abstract} \\end{center}  We study the strong gravitational lensing on the equatorial plane of a quasi-Kerr compact object with arbitrary quadrupole moments which can be used to model the super-massive central object of the galaxy. We find that, when the quadrupolar correction parameter $\\xi$ takes the positive (negative) value, the photon-sphere radius $r_{ps}$, the minimum impact parameter $u_{ps}$, the coefficient $\\bar{b}$, the relative magnitudes $r_m$ and the angular position of the relativistic images $\\theta_{\\infty}$ are larger (smaller)  than the results obtained in the Kerr black hole, but the coefficient $\\bar{a}$, the deflection angle $\\alpha(\\theta)$ and the angular separation $s$ are smaller (larger) than that in the Kerr black hole.    These features may offer a way to probe special properties for some rotating compact objects by the astronomical instruments in the future. ",
        "introduction": "In the framework of general relativity, the no-hair theorem \\cite{noh} guarantees that a neutral rotating astrophysical black hole is uniquely described by the Kerr metric only with two parameters, the mass $M$ and the rotational parameter $a$. Observations of the weak gravitational systems agree well with the prediction of the general relativity. However, the hypothesis that the astrophysical black-hole candidates are described by the Kerr metric still lacks the direct evidence, and the general relativity has been tested only for weak gravitational fields. In the regime of strong gravity, the general relativity could be broken down and astrophysical black holes might not be the Kerr black holes as predicted by the no-hair theorem \\cite{FCa,TJo,CBa,psa08}. Several parametric deviations from the Kerr metric, i.e., multipole moments, have been suggested to study observational signatures  including gravitational waves from extreme mass-ratio inspiral (EMRI) \\cite{gwave,ga1,ga2,JGa,Apostolatos} and the electromagnetic spectrum emitted by the accreting disk around black holes \\cite{ag1,CBa3}.  The general stationary axisymmetric  neutral compact object can been described in terms of its mass, rotational parameter and multipole moments \\cite{moments}. The multipole moments are consist of a set of mass multipole moment $M_l$ and current multipole moment $S_l$, here the subscript $l$ of them is labeled by the angular inter eigenvalue $l\\geq0$. The relation between the parameters of the multipole moments can be expressed as \\cite{moments} \\begin{eqnarray} M_l + {\\rm i}S_l = M({\\rm i}a)^l + \\delta M_l + {\\rm i}\\delta S_l. \\label{deltamult} \\end{eqnarray} For the Kerr black hole, the deviation $\\delta M_l$ and $\\delta S_l$ are equal to zero. Thus the relation (\\ref{deltamult}) is unique to the Kerr black hole and is just a mathematical expression of the famous no-hair theorem. Therefore, for a general rotating spacetime, by measuring three independent multipole moments of the spacetime we could provide sufficient information for verifying whether or not the central body is a Kerr black hole. Now, four approaches for measuring the multipole moments or a perturbation of the Kerr black hole have been proposed to test the no-hair theorem. The first approach based on writing the general stationary axisymmetric metric in terms of multipole moments was proposed by Ryan and was extended by Barack \\textit{et al}.\\cite{gwave}. They demonstrated how the laser interferometer space antenna (LISA) detector could map compact object deviating from the Kerr metric by means of EMRI observations. In the second approach, Collins and Hughes \\cite{ga1} introduced ``bumpy black hole\" (a spacetime that slightly deviates from the exact black hole of general relativity). This approach was recently generalized to the case of a rotating black hole in an alternate theories of gravity \\cite{ga2,ag4}. In the third one,  based on the Manko-Novikov metric \\cite{ag3}, Apostolatos \\textit{et al} \\cite{Apostolatos} showed that the appearance of Birkhoff chains in the neighborhood of a resonant tori would lead to such a modification of a gravitational-wave measurement. Finally, in the fourth one, Glampedakis \\textit{et al} \\cite{Glampedakis} proposed a quasi-Kerr metric by choosing the (dimensionless) quadrupole moment to be $q_{\\rm Kerr}-\\xi$, where the quadrupolar correction  parameter $\\xi$  represents a potential deviation from the Kerr metric and $q_{\\rm Kerr}=-J^2/M^4=-a^2/M^2$. And the quadrupole moment \\cite{Glampedakis} can be written as \\begin{equation} Q=-M\\left(a^2+\\xi M^2\\right). \\label{qradmoment} \\end{equation} Glampedakis \\textit{et al} calculated the periastron precession and constructed `kludge' gravitational waveforms as a function of the parameter $\\xi$. These waveforms can be significantly different from the expected Kerr signal even for small changes of the quadrupole moment. In a series of papers \\cite{Johannsen}, Johannsen \\textit{et al} analyzed in detail the innermost stable circular orbit, the circular photon orbit and the electromagnetic spectrum of the quasi-Kerr spacetime to test the no-hair theorem.  Gravitational lensing caused by deflection of light rays in a gravitational field is an ordinary phenomenon in astronomical observations when the light pass through the massive compact objects such as black hole, quasars and supernova.  In the last decades, the theory of gravitational lensing has been developed along two branches. The former takes the weak deflection limit as photon radius is much larger than the gravitational radius of the lens. The latter uses the strong deflection limit as the light ray loops around the massive compact object many times before it reaches to the observer. By this mechanism, two infinite series of relativistic images appears on each side of the lens. These relativistic images can provide us not only some important signatures about compact objects in the universe, but also profound verification of alternative theories of gravity in the strong field regime   \\cite{Vir,Vir1,Vir2,Vir3, Bozza2,Bozza3}. Thus, the strong gravitational lensing is regarded as a powerful indicator of the physical nature of the central celestial objects and then has been studied extensively in various theories of gravity \\cite{Gyulchev,Gyulchev1,Fritt,Bozza1,Eirc1,whisk,Bhad1,Song1,Song2,TSa1,AnAv, Ls1,Darwin}.   The main purpose of this paper is to study the strong gravitational lensing by the quasi-Kerr compact object \\cite{Glampedakis} and to see whether it can leave us the signature of the quadrupole moment parameter in the photon sphere radius, the deflection angle, the coefficients and the observable quantities of strong gravitational lensing. Moreover, we will explore how it differs from the Kerr black hole lensing.  The paper is organized as follows: In Sec. II, we will review briefly the metric of the quasi-Kerr compact object with the quadrupole moment parameter proposed by Glampedakis \\textit{et al} \\cite{Glampedakis}  and calculate the radius of photon sphere. In Sec. III, we study the physical properties of the strong gravitational lensing by  quasi-Kerr compact object and probe the effects of the quadrupole moment parameter on the deflection angle, the coefficients and the observable quantities for gravitational lensing in the strong field limit. We end the paper with a summary. ",
        "conclusions": "In this paper, in order to test how astronomical compact object deviates from the Kerr black hole, we investigate the features of light propagation on the equatorial plane of the rotating quasi-Kerr spacetime proposed by Glampedakis \\textit{et al} \\cite{Glampedakis}. Assuming that the massive compact object at the center of our galaxy can be described by quasi-Kerr spacetime, we obtain the photon radius, the coefficients and observable quantities of the strong gravitational lensing. We find that the photon-sphere radius exist only when the quadrupolar correction parameter takes the value $\\xi>-0.302$ for the photons in the prograde orbit $(a>0)$. Moreover, as the quadrupolar correction parameter $\\xi$ becomes negative the radius of the photon sphere becomes smaller,  which implies that the photons are more easily captured by the quasi-Kerr compact object with the negative quadrupolar correction parameter $\\xi$ than that of the Kerr black hole. It is interesting to note that, when the quadrupolar correction parameter $\\xi$ takes the positive (negative) value, the photon-sphere radius $r_{ps}$, the minimum impact parameter $u_{ps}$, the coefficient $\\bar{b}$, the relative magnitudes $r_m$ and the angular position of the relativistic images $\\theta_{\\infty}$ are larger (smaller)  than the results obtained in the Kerr black hole, but the coefficient $\\bar{a}$, the deflection angle $\\alpha(\\theta)$ and the angular separation $s$ are smaller (larger) than that of the Kerr black hole. Based on the above results, we come to the conclusion that there are some significant effects of the quadrupolar correction parameter $\\xi$ on the coefficients and observable parameters of the strong gravitational lensing.  These results, in principle, may provide a possibility to test how astronomical black holes with arbitrary quadrupole moments deviate from the Kerr black hole in the future astronomical observations."
    },
    "1208/1208.4657_arXiv.txt": {
        "abstract": "Most star formation in our galaxy occurs within embedded clusters, and these background environments can affect the star and planet formation processes occurring within them.  In turn, young stellar members can shape the background environment and thereby provide a feedback mechanism. This work explores one aspect of stellar feedback by quantifying the background X-ray radiation fields produced by young stellar objects.  Specifically, the distributions of X-ray luminosities and X-ray fluxes produced by cluster environments are constructed as a function of cluster membership size $N$. Composite flux distributions, for given distributions of cluster sizes $N$, are also constructed. The resulting distributions are wide and the X-ray radiation fields are moderately intense, with the expected flux levels exceeding the cosmic and galactic X-ray backgrounds by factors of $\\sim10-1000$ (for energies 0.2 -- 15 keV).  For circumstellar disks that are geometrically thin and optically thick, the X-ray flux from the background cluster dominates that provided by a typical central star in the outer disk where $r \\ga 9 - 14$ AU.  In addition, the expectation value of the ionization rate provided by the cluster X-ray background is $\\zeta_X\\sim8\\times10^{-17}$ s$^{-1}$, about 4 -- 8 times larger than the canonical value of the ionization rate from cosmic rays.  These elevated flux levels in clusters indicate that X-rays can affect ionization, chemistry, and heating in circumstellar disks and in the material between young stellar objects. ",
        "introduction": "Most star formation in our galaxy occurs within embedded clusters, which are themselves located inside giant molecular clouds.  The radiation fields produced by these stellar nurseries can influence the formation of additional cluster members, and especially their accompanying planetary systems (Adams 2010). More specifically, the processes of star and planet formation can be affected via [1] heating of starless cores, leading to evaporation and the loss of star forming potential (e.g., Gorti \\& Hollenbach 2002), [2] ionization within starless cores, leading to greater coupling between the magnetic fields and gas (e.g., Shu 1992), thereby acting to suppress continued star formation, [3] evaporation of circumstellar disks, leading to loss of planet forming potential (e.g., Shu et al. 1993, Hollenbach et al. 1994, St{\\\"o}rzer \\& Hollenbach 1999, Adams et al. 2004), and [4] ionization of circumstellar disks, which helps maintain the magneto-rotational instability (MRI), which in turn helps drive disk accretion (e.g., Balbus \\& Hawley 1991).  With regard to circumstellar disks, we note that the background radiation from the cluster environment often dominates that produced by the central star, especially at ultraviolet (UV) wavelengths (e.g., Hollenbach et al. 2000; Armitage 2000; Adams \\& Myers 2001, Fatuzzo \\& Adams 2008); nonetheless, the photoevaporation of disks is often driven by stellar radiation (e.g., Shu et al. 1993; Alexander et al. 2004, 2005, 2006).  Multi-wavelength observations of embedded clusters over the past two decades have provided a wealth of information on these environments (Lada \\& Lada 2003; Porras et al. 2003; Allen et al. 2007), making it possible to assess what impact a given physical process can have on star and planet formation. Indeed, a comprehensive analysis of the effects of FUV and EUV background fields on cluster environments has already been performed (Armitage 2000; Adams et al. 2006; Fatuzzo \\& Adams 2008; Hollenbach \\& Gorti 2009; Holden et al. 2010). The main goal of this paper is to perform a similar analysis for the X-ray band.  This paper is organized as follows. In Section 2, we discuss the cluster environment, present our characterization of the stellar IMF, and outline the basic approach for assigning X-ray luminosities to stars in our cluster models.  We discuss the statistical aspects of our cluster model in Section 3 and present distributions of the X-ray luminosities for clusters with varying stellar membership size $N$. In Section 4, we construct the corresponding distributions of X-ray fluxes, both for varying $N$ and for different forms for the stellar density profiles. The implications of these findings are discussed in Section 5; the paper concludes in Section 6 with a summary of results and a discussion of potential applications. ",
        "conclusions": "This paper constructs the expected X-ray radiation fields provided by young embedded clusters. Specifically, we calculate the distributions of X-ray luminosity for clusters of a given membership size $N$ and the corresponding distributions of X-ray flux. We also determine the distributions of X-ray flux for two cluster ensembles (for the cluster distribution found in the solar neighborhood and for an extended distribution with maximum membership size $N_{max}=10^5$).  The flux distributions depend on the density profile of the stellar objects (the individual X-ray sources) within the cluster and we consider models that span the expected range of possibilities. The main results can be summarized as follows:  The expected flux levels for X-rays are modest, with a characteristic value of order $F_X \\sim 1-6 \\times 10^{-5}$ erg cm$^{-2}$ s$^{-1}$; the range corresponds to the types of clusters under consideration, where the X-ray flux increases with cluster membership size $N$ and with the central concentration of the density profile (Figures \\ref{fig:fluxdist} -- \\ref{fig:flux_compbig}).  The distributions of X-ray flux are relatively wide, however, with the full width at half-maximum for $\\log_{10} F_X$ corresponding to a factor of $\\sim3-4$ for $\\rho\\propto 1/r$ density profiles and a factor of $\\sim5-10$ for $\\rho\\propto 1/r^2$ density profiles (see Figures \\ref{fig:fluxdist} and \\ref{fig:fluxdistalt}). The cluster to cluster variation in X-ray flux will thus be significant.  For completeness, we also note that X-ray emission is highly variable with time, and that the scatter in X-ray luminosity generally decreases with cluster age (Alexander \\& Preibisch 2012).  Given the expected X-ray flux distributions for embedded clusters, as calculated herein, the outer regions of circumstellar disks (where planets form) receive comparable amounts of flux from their central stars and from the background. The central star dominates for small clusters if one ignores geometric effects (equation [\\ref{fxest}]), whereas the background X-ray flux from the cluster dominates for geometrically thin disks (equation [\\ref{xraystar}]) and also for larger clusters (Figure \\ref{fig:flux_compbig}).  Circumstellar disks are expected to be relatively thin and the flux impinging on them from the central star is significantly reduced by geometric effects. As a result, the background X-ray flux is expected to exceed that of the star for disk radii beyond $r\\approx 9-14$ AU. We stress that significant variability in these values are expected given both the large variability in the X-ray luminosities of individual stars and the broad X-ray flux distributions within the cluster environment.  Whereas disks receive comparable X-ray fluxes from their central stars and the background cluster, the situation is markedly different for the case of both FUV and EUV radiation. At UV wavelengths, the background radiation fields of the cluster overwhelm those of the central stars for clusters with $N \\ge 100$ (Armitage 2000; Adams \\& Myers 2001; Fatuzzo \\& Adams 2008). As a result, photoevaporation of circumstellar disks due to UV radiation is often dominated by the cluster background (not the central star), whereas mass loss due to X-rays can be dominated by the central star (and not the cluster).  The ionization rate provided by X-rays in the cluster environment has a characteristic value $\\zeta_X \\sim 8 \\times 10^{-17}$ s$^{-1}$ (see equation [\\ref{zetaxnumber}]). This ionization rate is about 4 -- 8 times the fiducial ionization rate $\\zeta_{CR}$ from cosmic rays, which are often considered to be the dominant source of ionization for molecular clouds and star formation (e.g., Shu 1992). The ionization rates for both X-rays $\\zeta_X$ and cosmic rays $\\zeta_{CR}$ have wide distributions, so that the relative importance of the two sources will vary significantly from cluster to cluster (and can vary with time). In a similar vein, the background of X-ray radiation produced by the cluster is larger than the background radiation provided by the galaxy or the universe. Over the energy range under consideration, the cluster contribution dominates by a factor in the range 10 -- 1000, with a typical value of $\\sim100$.  The X-ray background radiation in young embedded clusters has a number of potential effects on star and planet formation (Glassgold et al. 2000; Feigelson et al.  2007).  The elevated levels of ionization lead to greater coupling between the cluster gas and magnetic fields (e.g., Shu 1992); this increased coupling, in turn, leads to longer time scales for ambipolar diffusion and acts to slow down additional star formation.  Enhanced ionization also affects MHD-driven turbulence, both in the cluster gas and in circumstellar disks via MRI (Balbus \\& Hawley 1991).  The background X-ray flux will produce line diagnostics in circumstellar disks (e.g., Tsujimoto et al. 2005; Hollenbach \\& Gorti 2009) and will affect chemical reactions in the gas (Aikawa \\& Herbst 1999). These processes, and many others, should be studied in the future to provide a more complete picture of how cluster environments shape the solar systems forming within them.  \\vskip 0.35truein \\centerline{\\bf Acknowledgment} We thank the referee for their careful review of the manuscript, and for their useful comments that improved the quality of the writing. FCA is supported at UM by NASA grant NNX11AK87G from the Origins of Solar Systems Program, and by NSF grant DMS-0806756 from the Division of Applied Mathematics. MF is supported at XU through the Hauck Foundation. LH is supported at NKU through a CINSAM Research Grant.  \\bigskip"
    },
    "1208/1208.6470_arXiv.txt": {
        "abstract": "The stationary axisymmetric force-free magnetosphere of a pulsar is studied analytically. The pulsar equation is solved in the region close to the magnetic axis. Proceeding from linearization of the current function in the axial region, we find the axial magnetic flux function valid at any altitude above the neutron star. This function is used as a starting approximation to develop series for the non-linear pulsar equation in the polar region. Taking into account the quasi-monopolar character of the pulsar magnetic flux at infinity, we obtain unique asymptotic series for the flux and current functions. At infinity, both functions are close but not equivalent to those known for the case of a force-free monopole. The flux function at the top of the polar gap is found to differ from the dipolar one at the neutron star surface. With our results, the transverse current sheet closing the pulsar circuit at the neutron star surface is consistently incorporated into the global magnetospheric structure, the backward particle flow at small polar angles can be excluded and the stationary cascade scenario looks admissible. The present paper is the first step toward complete analytic description of the pulsar force-free magnetosphere allowing for the plasma-producing gaps and pulsar current circuit closure. ",
        "introduction": "Pulsar magnetosphere contains the plasma \\citep{gj69}, which is sufficiently abundant to affect the magnetic field structure. This necessitates a self-consistent consideration of fields and currents in pulsars. The basic model of the pulsar magnetosphere is that of a rotating axisymmetric force-free dipole, where the magnetic and rotational axes are aligned, the electromagnetic forces are balanced and the particle inertia is ignored. Then the poloidal current and magnetic flux are related by the well-known pulsar equation \\citep{m73,sw73,o74}. Being formulated almost 40 years ago, it still lacks a proper analytic solution for the dipolar case.  An exact solution is known only for the force-free monopole \\citep{m73}, including the versions of a split and offset monopole \\citep{m91,p12}, which are believed to replicate the behaviour of the force-free dipole at large distances from the origin. In the case of a monopole, the problem is strongly simplified, since the poloidal component of the force-free magnetic field appears to be the same as in the original vacuum configurations. In the case of a dipole, the flux and current functions entering the pulsar equation are both unknown, and the problem is to guess such a current distribution that makes the self-consistent flux function obey the physically meaningful boundary conditions. The early attempts of solving the pulsar equation for the current functions of a special form yielded the flux functions which were valid only inside the light cylinder and could not be smoothly continued to infinity \\citep*{m73b,bgi83}. Later on a more sophisticated current function providing a smooth passage of the magnetic field lines across the light cylinder was found by means of numerical methods \\citep*{ckf99}. The necessity to revise the set of commonly used boundary conditions was pointed out in \\citet{p12}.  The pioneering work of \\citet{ckf99} inspired extensive numerical studies of the pulsar equation. The numerical solution for the stationary axisymmetric force-free magnetosphere was confirmed by the time-dependent simulations \\citep{k06,mck06,s06} and generalised in a number of aspects. In particular, the non-axisymmetric 3D magnetosphere was considered \\citep*{s06,kc09,bs10,kck12}, the differential rotation was taken into account \\citep{c05,t06,t07}, the non-ideal MHD case was addressed \\citep*{kkhc12,lst12}. However, the set of boundary conditions used in all these simulations did not allow for the presence of the plasma producing gaps. Therefore the numerical treatment of the pulsar equation was not self-consistent.  In our previous work \\citep{p12}, we have suggested an empirical model of the pulsar axisymmetric force-free magnetosphere, which for the first time includes the polar, outer and slot gaps into the global magnetospheric structure. The basic prescriptions of the model are as follows. Although the polar gap height can be regarded as infinitesimal, the magnetic field structure at its top differs from the original dipolar one at the neutron star surface. As the outer gap is entirely located within the light cylinder, beyond the light cylinder the null line should not intersect the magnetic field lines. The polar and outer gaps control different bundles of the open field lines and produce the direct and return currents, respectively. Therefore beyond the light cylinder the null line coincides with the critical field line, both going parallel to the magnetic equator at a certain altitude above it. The compensation current necessary to close the pulsar circuit originates in the slot gap located along the critical field line inside the light cylinder.  Our qualitative model needs to be quantified. With the present paper, we begin a systematic analytic description of the pulsar magnetosphere within the framework of this model. The present study is focused on the region close to the magnetic axis. The absence of the poloidal current along the magnetic axis provides an essential simplification: in the nearest vicinity of the axis, the generally non-linear current function and the pulsar equation on the whole can be linearised. In Sect.~2, we solve the simplified pulsar equation in the axial region and find the magnetic flux function at any altitude $z$ above the neutron star. In Sect.~3, the complete non-linear pulsar equation is considered. We search for the solution in terms of series in $\\rho^2/z^2$  (where $\\rho$ is the separation from the axis) and use the axial flux function as a starting approximation. Taking into account the quasi-monopolar character of the flux function at infinity, we find a unique current function and the corresponding unique asymptotic form of the flux function at large altitudes above the neutron star. With the current function in hand, we also obtain the flux function at the top of the polar gap and demonstrate its distinction from that of a dipole. Our results and their implications are discussed in Sect.~4 and briefly summarised in Sect.~5. ",
        "conclusions": "\\protect\\label{s5}  We have considered the stationary axisymmetric force-free magnetosphere of a pulsar and performed the analytic treatment of the pulsar equation close to the magnetic axis. Linearization of the current function in the axial region leads to the self-consistent flux function with the approximately dipolar and monopolar behaviour at low and high altitudes, respectively. The axial flux function has been used to construct the solution of the non-linear pulsar equation at small polar angles in terms of series. Keeping in mind the quasi-monopolar character of the flux function at infinity, we have obtained unique asymptotic series for the flux and current functions, which differ from the monopolar ones as $f-f_\\mathrm{mon}\\sim -\\rho^6/8z^6$, $AA^\\prime-AA^\\prime_\\mathrm{mon}\\sim -4f^3$. With the current function obtained, the flux function at the top of the polar gap appears to differ from the dipolar one as $f-f_\\mathrm{dip}\\sim 7\\rho^8/96z^9$. This is attributed to the action of the transverse current closing the current circuit. Thus, our result includes the polar part of the transverse current sheet at the neutron star surface into the global magnetospheric structure.  Possible implications of our results can be summarized as follows. \\begin{enumerate} \\item The axial flux function found implies an additional boundary condition at the axis, $(\\partial^2f/\\partial\\rho^2)_{\\rho=0}=2\\Phi(z)$. This may, at least partially, compensate for the unknown boundary segments and boundary conditions in the pulsar force-free problem. \\item The self-consistent magnetic field at the top of the polar gap may constrain the physics of current closure in pulsars and give new insights into the polar gap theory. \\item Our results testify against the backward particle flow at small polar angles. \\item The relation of the resultant net charge and poloidal current densities does not exclude the stationary cascade scenario in the polar region of the polar gap. \\end{enumerate}  The results of the present paper are the first step in the analytic description of our model of the pulsar force-free magnetosphere \\citep{p12}, which consistently includes the plasma-producing gaps. As our model assumes the same axial boundary conditions as the classical one, in the polar region considered here our results are roughly in line with the numerical results obtained within the framework of the classical model. However, in other parts of the magnetosphere our model is essentially distinct, and the analytic description of fields and currents in these regions is of particular interest. This will be a subject of the subsequent papers."
    },
    "1208/1208.3564.txt": {
        "abstract": "This is the third paper on the improvements of systematic errors in our weak lensing analysis using an elliptical weight function, called E-HOLICs. In the previous papers we have succeeded in avoiding error which depends on ellipticity of background image. In this paper, we investigate the systematic error which depends on signal to noise ratio of background image. We find that the origin of the error is the random count noise which comes from Poisson noise of sky counts. Random count noise makes additional moments and centroid shift error, and those 1st orders are canceled in averaging, but 2nd orders are not canceled. We derived the equations which corrects these effects in measuring moments and ellipticity of the image and test their validity using simulation image. We find that the systematic error becomes less than 1\\% in the measured ellipticity for objects with $S/N>3$. ",
        "introduction": "The importance of the weak lensing analysis is now widely recognized because it has a potential to provides us a direct and unbiased information on the mass distribution for lens objects. The weak lensing analysis measures shapes(called ellipticity which has two components interpreted as direction and magnitude) of many background images(galaxies) and then averaged over an appropriate number of images to get rid of intrinsic random ellipticity of images and to withdraw the ellipticity due to gravitational tidal effect(shear) of the lensing object. The shear carries the information of mass structure of the lensing object. Thus an accurate shape measurement of the background images is critically important to accurately measure the mass distribution. So far weak lensing is very successful for cluster lensing (ellipticity due to shear is of the order of 5\\%) and provides us a rich information of mass structures of clusters and of our understanding structure formation in the universe. Recently the cosmic shear, i.e. the weak lensing due to large scale structure(LSS) attracted much attention because of it's ability to study the nature of dark energy which is supposed to be the source of the accelerated expansion of the universe. In fact several projects for the cosmic shear measurement are proposed and some of them is almost ready to start the observation (Hyper Suprime-Cam http://www.naoj.org/Projects/HSC/HSCProject.html, Dark Energy Survey http://www.darkenergysurvey.org/, Euclid http://sci.esa.int/euclid and so on). However the signal of cosmic shear is very weak(of the order of 1\\%) compared with cluster lensing and thus needs special treatment.  Namely we needs to develop very accurate shape measurement scheme which avoids various systematic errors. For example the measured gravitational shear depends on the ellipticity and signal to noise of background image. Usually such dependence becomes small by averaging many of the images, but it is critically important to realize that these dependence somehow correlated with the redshift distribution of image which is also important to have an accurate measurement of the shear. Thus we cannot make a simple averaging over the images without having a method free from such systematic errors. The required accuracy for the measurement of ellipticity is less than 1\\% in order to have an useful information of dark energy.  There have been many studies in this direction and various measurement schemes are proposed(Kaiser et al 1995, Bernstein \\& Jarvis 2002; Refregier 2003; Kuijken et al. 2006; Miller et al. 2007; Kitching et al. 2008; Melchior 2011). The accuracy of these methods are tested using the simulation data provided by STEP1(Heymans et al 2006), STEP2(Massey et al 2007), GREAT08(Bridle et al 2010) and GREAT10(Kitching et al 2012). Although much progress is reported, none of the method achieved the required accuracy and are free from various systematic errors.  We have also developed a new scheme based on KSB method(Kaiser et al 1995) using an elliptical window function(we called E-HOLICs) to measure the background image as accurately as possible(Okura and Futamase 2011, Part I paper). It is shown in our Part II papers that the E-HOLICs can improve the systematic error which depends on ellipticity. In this paper we study the systematic error which depends on signal to noise ratio(SN). There are some studies about this systematic error (Hirata et al 2004, Kacprzak et al 2012, Refregier et al 2012, Okura and Futamase 2012 Part II paper, Melchior and Viola 2012). these results show this systematic error comes from random count noise(RCN). Because, 1st order effects from RCN are canceled by averaging, but 2nd order effects are not canceled. We calculate the 2nd order effects to obtain the correction formulas in the measurement of moments and ellipticity  for Gaussian weighted images in KSB method and E-HOLICs method (i.e. without PSF correction). Using the simulation data GREAT08, we find that the derived formula correct the SN dependent bias within 1\\% for images with SN $\\geq$ 3.  The paper organized as follows. In section 2, we explain and define our notations and some of the definitions. In section 3, we calculate the 2nd order effects of RCN and obtain general formulas to correct the effects. We test the formula in the case of KSB method with Gaussian weight function. The correction formula in the case of E-HOLICs is presented in section 4, and tested it using GREAT 08 simulation and find that the systematic error becomes less than 1\\% in the measured ellipticity for objects with $S/N>3$. In section 5, we summaries our results. %---------------------------------------------------------------------------------------------------- ",
        "conclusions": "Following the previous work we studied in this paper the systematic error caused by signal to noise(SN) ratio of the observed image in our weak lensing analysis called E-HOLICs. It has been known that the shear is underestimated when low SN background images are used and is overestimated when high SN background images are used in the weak lensimng analysis. The latter error was improved in the previous work. The improvement of the former error is important because if we have such improvement, we will be able to use many faint background sources which improves the statistical accuracy of the weak lensing analysis.  We have identified the origin of the systematic error as the photon random count noise by sky. Although its 1st order effect vanishes by averaging, but 2nd order effects are not canceled in measuring the moments and centroid of the images. We investigated this effect carefully and obtain the formulas in KSB method and E-HOLICs method to correct the effect in measuring moments and ellipticity. Although general expressions for these formulas are complicated, they reduce to relatively simple forms for images with an elliptical Gaussian form(EGI). We tested the validity of the correction formula eq.(\\ref{eq:cformH}) and eq.(\\ref{eq:cformD}) for EGIs using simulatoin data. Furthermore we applied the general formulas to GREAT08 and confirmed that the systematic error reduces to less than 1\\% in measuring ellipticity for images with WSN=9.1 which roughly corresponds to SN=3 object.  Although the present analysis has not taken into account the Point Spread function(PSF) correction which is necessary for the observation from the ground. PSF correction uses complicated combinations of higher moments and will be very complicated in E-HOLICs approach. However the above result is very encouraging and is worthwhile challenging. Finally we should point out that the present work will be applicable for the space based observation because PSF by instrument is expected to be small for such observation. It will be very interesting to confirm this expectation by using data such as COSMOS.  We thank Satoshi Miyazaki, Takashi hamana, Keiichi Umetsu,  Nobuhiro Okabe, Yousuke Utsumi and Yuichi Higuchi for useful discussions and comments. This work is partially supported by the COE program''Weaving Science Web beyond Particle-matter Hierarchy'' at Tohoku University and  Grant-in -Aid for Scientific Research from JSPS(Nos.18072001, 23540282 for TF). %---------------------------------------------------------------------------------------------------- %---------------------------------------------------------------------------------------------------- \\appendix"
    },
    "1208/1208.5897.txt": {
        "abstract": "%% Text of abstract Cosmic rays are a sample of solar, galactic and extragalactic matter. Their origin and properties are one of the most intriguing question in modern astrophysics. The most energetic events and active objects in the Universe: supernovae explosion, pulsars, relativistic jets, active galactic nuclei, have been proposed as sources of cosmic rays although unambiguous evidences have still to be found. Electrons, while comprising $\\sim 1\\%$ of the cosmic radiation, have unique features providing important information regarding the origin and propagation of cosmic rays in the Galaxy that is not accessible from the study of the cosmic-ray nuclear components due to their differing energy-loss processes. In this paper we will analyse, discussing the experimental uncertainties and challenges, the most recent measurements on cosmic-ray nuclei and, in particular, electrons with energies from tens of GeV into the TeV region. ",
        "introduction": "We would like to thank the PAMELA Collaboration for providing some of the information included in this paper and E. Mocchiutti for helping in the production of several figures.  % ",
        "conclusions": ""
    },
    "1208/1208.0126_arXiv.txt": {
        "abstract": "\\label{abstract}  The ongoing characterization of hot Jupiters has motivated a variety of circulation models of their atmospheres.  Such models must be integrated starting from an assumed initial state, which is typically taken to be a wind-free, rest state.  Here, we investigate the sensitivity of hot-Jupiter atmospheric circulation models to initial conditions.  We consider two classes of models---shallow-water models, which have proven successful at illuminating the dynamical mechanisms at play on these planets, and full three-dimensional models similar to those being explored in the literature.  Models are initialized with zonal jets, and we explore a variety of different initial jet profiles.  We demonstrate that, in both classes of models, the final, equilibrated state is independent of initial condition---as long as frictional drag near the bottom of the domain and/or interaction with a specified planetary interior are included so that the atmosphere can adjust angular momentum over time relative to the interior.  When such mechanisms are included, otherwise identical models initialized with vastly different initial conditions all converge to the same statistical steady state.  In some cases, the models exhibit modest time variability; this variability results in random fluctuations about the statistical steady state, but we emphasize that, even in these cases, the statistical steady state itself does not depend on initial conditions.  Although the outcome of hot-Jupiter circulation models depend on details of the radiative forcing and frictional drag, aspects of which remain uncertain, we conclude that the specification of initial conditions is not a source of uncertainty, at least over the parameter range explored in most current models. ",
        "introduction": "\\label{Introduction} Since the discovery of the first exoplanet around a main-sequence star, 51~Pegasi~b \\citep{mayor-queloz-1995}, almost 800 planets orbiting other stars have been detected.  Nearly $20\\%$ of them are hot Jupiters \\citep{wright-etal-2011}, giant planets orbiting within $\\sim$$0.1\\,$AU of their central stars.  The physical regime of hot Jupiters differs substantially from those of solar-system giant planets; the presumed tidal locking leads to modest rotation rates and permanent day and nightsides, with incident stellar fluxes $\\sim$$10^3$--$10^6$ times stronger than that received by giant planets in our solar system.  These differences have motivated a flourishing research program focused on elucidating the atmospheric dynamics of these worlds.  Current observations place meaningful constraints on planetary radii, atmospheric composition, albedo, and the three-dimensional temperature structure, including dayside temperature profiles and variation of the temperature between day and night \\citep[e.g.,][]{knutson-etal-2007b,charbonneau-etal-2008, knutson-etal-2008, cowan-etal-2007, harrington-etal-2006, harrington-etal-2007}.  These observations have motivated a variety of three-dimensional (3D) circulation models of hot Jupiters \\citep{showman-guillot-2002, cooper-showman-2005, cooper-showman-2006, showman-etal-2008a, showman-etal-2009, menou-rauscher-2009, rauscher-menou-2010, rauscher-menou-2012, dobbs-dixon-lin-2008, dobbs-dixon-etal-2010, lewis-etal-2010, perna-etal-2010, perna-etal-2012, thrastarson-cho-2010, thrastarson-cho-2011, heng-etal-2011, heng-etal-2011b, showman-etal-2012}.  These models obtain a circulation pattern generally comprising several broad atmospheric jets, including in most cases a strong eastward equatorial jet (superrotation) with speeds up to a few $\\rm km\\,s^{-1}$ and westward zonal-mean flow at high latitude.  An analytical explanation was provided by \\cite{showman-polvani-2011}, who showed that the superrotation results from the interaction of planetary-scale Rossby and Kelvin waves---themselves a response to the day-night thermal forcing---with the mean flow.  When the radiative and advection time scales are comparable, the hottest region can be displaced eastward from the substellar point by tens of degrees longitude \\citep[e.g.,][]{showman-guillot-2002}, as subsequently observed on HD 189733b \\citep{knutson-etal-2007b}.  Nevertheless, it is notable that some 3D models produce qualitatively different circulation patterns exhibiting a range of flow behavior \\citep{thrastarson-cho-2010}.  An important issue in modeling atmospheric circulation is the choice of initial conditions.  Most of the models published to date integrate the equations starting from a rest state containing no winds, although \\citet{cooper-showman-2005} also performed an integration starting from an initial condition containing a broad westward jet and found that the initial condition did not strongly affect the outcome. Nevertheless, \\citet{thrastarson-cho-2010} recently performed a detailed investigation and found that, in their model, the initial conditions severely affect the final state.  They explored initial conditions containing a broad eastward jet, broad westward jet, and three-jet pattern, all with speeds of 0.5--$1\\rm\\,km\\,s^{-1}$, and compared this to models integrated from rest.  These models---differing only in the initial condition---led to a wide range of final states whose pattern of flow streamlines and the horizontal temperature distribution (including the longitudinal offsets of any hot or cold spots) differed significantly. \\citet{thrastarson-cho-2010} found that such extreme sensitivity to initial condition occurred even in models whose initial conditions differed only slightly.     From a mathematical perspective, the initial conditions of course comprise an essential aspect of defining the overall mathematical problem, and in many dynamical systems, the initial conditions indeed play a crucial role in affecting the solution.  However, atmospheres are forced-dissipative systems, and such forcing and dissipation tend to drive the atmospheric circulation into a statistical steady state retaining little if any memory of its initial condition.  For this reason, many terrestrial general circulation model (GCMs) investigating the statistical steady state of Earth's global atmospheric circulation are integrated from rest,\\footnote{This of course is not true for weather prediction models, which consist of very-short-term integrations of typically a few days, determining how some given observed circulation evolves with time.} with no expectation that this choice adversely influences the results.  A possible exception to this expected lack of sensitivity would be if the atmospheric circulation exhibits multiple, stable equilibria corresponding to radically different circulation patterns for an identical set of forcing and boundary conditions.  In this case, determining which of these multiple equilibria the atmosphere resides in requires knowledge of its history, potentially including its initial condition.  An example in the context of terrestrial-planet climate is the existence of stable equilibria corresponding to primarily ice-free and globally glaciated (``Snowball Earth'') states \\citep{budyko-1969}.  In the range of incident stellar fluxes where both equilibria exist, the actual state occupied by the climate depends on history \\citep[for specific examples of how this works, see][]{pierrehumbert-2010}.  In just such a way, the question of initial-condition sensitivity for the atmospheric circulation is therefore perhaps best thought of as a question of whether the atmosphere exhibits multiple rather than only one stable equilibrium. To date, no claims have been made in the literature that the atmospheric circulation of hot Jupiters exhibit multiple, stable equilibria for a given set of forcing conditions, and so a sensitivity to initial conditions \\citep{thrastarson-cho-2010} is not expected {\\it a priori}.  Nevertheless, the issue is worthy of further study.   Here, we present the results of a thorough exploration of the initial condition sensitivity in atmospheric circulation models of hot Jupiters. We investigate two classes of models---idealized shallow-water models in Section~\\ref{2D model}, as these have recently been used to identify dynamical mechanisms operating in the atmospheres of hot Jupiters \\citep{showman-polvani-2011}---and fully 3D models in Section~\\ref{3D model}. Section~\\ref{conclusion} concludes. ",
        "conclusions": "\\label{conclusion}  We explored the sensitivity to initial conditions of three-dimensional models of synchronously rotating hot Jupiters with day-night thermal forcing.  The thermal forcing was chosen to be strong at the top and weak at the bottom, as must occur on real hot Jupiters.  Models were integrated from rest and from various eastward and westward jet profiles with speeds up to $\\sim$$1.5\\rm\\,km\\,s^{-1}$.  In all models explored, we found that the statistical steady states are independent of initial conditions---as long as the flow is anchored by interaction with a planetary interior so that the angular momentum of the atmosphere can adjust relative to that of the interior.  In the context of atmosphere models, such interaction could be parameterized by frictional drag near the bottom of the domain or by allowing the atmosphere to exchange mass, energy, and angular momentum with a specified abyssal layer underlying the atmosphere.  When such an interaction is included, all our models---for a given set of forcing parameters---converged to the same statistical steady state regardless of the initial condition employed.  When the thermal forcing is strong, the circulation in the equilibrated state exhibits modest time variability that induces small-amplitude random fluctuations in any given realization.  The statistical steady state itself, including not only the time-mean wind and temperature but the overall amplitude of these fluctuations, are independent of initial conditions when drag---or direct interaction with a specified abyssal layer---are included.  As described in the Section~\\ref{Introduction}, the issue of initial-condition sensitivity is perhaps best thought of in terms of whether the atmospheric circulation exhibits a single, rather than multiple, stable equilibria.  Taken at face value, our models suggest empirically that, for any given set of forcing and damping parameters, there exists only one stable equilibrium---at least for the range of forcing, damping, and planetary parameters explored here.  We have intentionally chosen forcing and planetary parameters similar to those appropriate to typical hot Jupiters, including HD 189733b and HD 209458b, as explored by a number of authors \\citep{showman-etal-2009, heng-etal-2011, heng-etal-2011b, perna-etal-2012, rauscher-menou-2012}. Therefore, we expect our fundamental result---the lack of sensitivity to initial conditions---to apply generally to the regimes explored in those papers.  In this context, it is interesting that our findings differ so drastically from those of \\citet{thrastarson-cho-2010}.  Their models differ from ours in two important ways: they lack large-scale drag, and the profiles of $T_{\\rm eq}$ and $\\tau_{\\rm rad}$ in their Newtonian-cooling scheme are independent of pressure.  Both these differences seem to contribute to the differences in their results relative to those presented here.  In particular, the absence of frictional drag in a 3D model with free-slip boundary conditions and no mass fluxes through the boundaries implies the absence of any external torques that could change the globally integrated angular momentum over time.  Thus, the globally integrated angular momentum of such a model will be conserved over time to within numerical accuracy. Since initial conditions corresponding to eastward jets, westward jets, and rest states exhibit different angular momenta, there is thus no mechanism to force those models to converge to the same angular momentum and hence final state.  In practice, we found that this sensitivity is not strong in the observable atmosphere for the range of initial conditions explored here.  Nevertheless, it may be stronger when $T_{\\rm eq}$ and $\\tau_{\\rm rad}$ are constant with depth, as is the case in most of \\citet{thrastarson-cho-2010}'s models.  Regardless, our results highlight the importance of anchoring the flow to an assumed planetary interior, either via the application of frictional drag; the introduction of a deep, quiescent layer at the bottom of the domain (as exists in many published 3D hot Jupiter models), or the explicit assumption of an abyssal layer underlying the active layer (as exists in 1-1/2 layer shallow-water models).  Only in this way can the atmosphere experience net torques that allow it to adjust angular momentum over time, allowing models with differing initial angular momenta to converge to a single statistical steady state.    \\vskip 5pt  Overall, our results indicate that specification of initial conditions is not a source of uncertainty in atmospheric circulation models of hot Jupiters, at least over the parameter range explored here.  This supports the continued use of hot-Jupiter GCMs for understanding dynamical mechanisms, explaining observations, and making predictions to help guide future observations.  That said, our results, as well as those of numerous previous publications, show that details of the radiative forcing and frictional damping significantly affect the flow structure, including the qualitative dynamical regime, wind speeds, day-night temperature differences, and longitudinal offsets of any hot or cold regions.    State-of-the-art GCMs now exist that include detailed non-grey radiative transfer \\citep{showman-etal-2009} as well as simpler, faster, gray treatments \\citep{heng-etal-2011b, rauscher-menou-2012, perna-etal-2012}.  By comparison, our understanding of how to specify frictional drag is less well developed, and areas such as inclusion of clouds, sub-gridscale parameterizations of turbulent mixing, and coupling to chemistry have received little attention.  Continued model development in these areas, and comparison of such models to observations, should improve our ability to discern the physical and dynamical regimes of these fascinating worlds.    \\vskip 20pt"
    },
    "1208/1208.5229_arXiv.txt": {
        "abstract": "We use both an HI-selected and an optically-selected galaxy sample to directly measure the abundance of galaxies as a function of their ``baryonic'' mass (stars + atomic gas). Stellar masses are calculated based on optical data from the Sloan Digital Sky Survey (SDSS) and atomic gas masses are calculated using atomic hydrogen (HI) emission line data from the Arecibo Legacy Fast ALFA (ALFALFA) survey. By using the technique of abundance matching, we combine the measured baryonic function (BMF) of galaxies with the dark matter halo mass function in a $\\Lambda$CDM universe, in order to determine the galactic baryon fraction as a function of host halo mass. We find that the baryon fraction of low-mass halos is much smaller than the cosmic value, even when atomic gas is taken into account. We find that the galactic baryon deficit increases monotonically with decreasing halo mass, in contrast with previous studies which suggested an approximately constant baryon fraction at the low-mass end. We argue that the observed baryon fractions of low mass halos cannot be explained by reionization heating alone, and that additional feedback mechanisms (e.g. supernova blowout) must be invoked. However, the outflow rates needed to reproduce our result are not easily accommodated in the standard picture of galaxy formation in a $\\Lambda$CDM universe. ",
        "introduction": "It is by now well established that baryonic matter represents only about $1/6$ of the total matter density of the universe \\citep[e.g.][]{Komatsu2011}, while the majority is in the form of non-baryonic dark matter (DM). Since galaxies form through the accretion of baryonic material onto dynamically dominant DM structures (halos), it would be reasonable to assume that the baryon mass fraction of present day galaxies approximately equals the cosmic value ($f_b = \\Omega_b / \\Omega_m \\approx 0.16$). Despite this expectation, observations point to the fact that galaxies are not able to retain their cosmic ``fair share'' of baryons, and that the resulting baryon deficit depends strongly on the mass of their host halo.  The first line of evidence is provided by observational estimates of the abundance of galaxies as a function of their total stellar mass, a distribution referred to as the galactic stellar mass function (SMF). Thanks to the advent of wide area optical surveys with multiband photometric and spectroscopic information, such as the Two degree Field Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey (SDSS), the SMF has been measured over the mass range $M_\\ast \\approx 10^7 - 10^{12} \\: M_\\odot$, using statistical samples of tens of thousands of galaxies and a variety of stellar mass estimation techniques (\\citealp{Cole2001}, \\citealp{Bell2003}, \\citealp{Panter2007}, \\citealp{Baldry2008}, \\citealp{LW2009}, \\citealp{Yang2009}, \\citealp{Baldry2012}, to name a few). The SMF displays an exponential cutoff at masses $M_\\ast \\gtrsim 10^{11} M_\\odot$ and an approximate  power-law behavior at low masses ($dn \\propto M_\\ast^{-\\alpha} \\: dM_\\ast $), with a ``shallow'' exponent of $\\alpha \\approx -1.3$. On the other hand, the halo mass function (HMF) predicted in the lambda cold dark matter ($\\Lambda$CDM) model, follows a much ``steeper'' power-law  ($\\alpha \\approx -1.8$) over the mass range of interest. This observation alone excludes the possibility that the stellar mass of a galaxy is simply a fixed fraction of the host halo mass.  One can furthermore statistically derive an average relation between the stellar mass of a galaxy ($M_\\ast$) and the mass of its host halo ($M_h$), through the technique of abundance matching (see \\S \\ref{sec:am} for details). $M_\\ast$ - $M_h$ relations based on abundance matching \\citep[e.g.][]{Guo2010, Moster2010, Behroozi2010, Leauthaud2012} have shown that the ``stellar conversion efficiency'', $\\eta_\\ast = (M_\\ast/M_h) \\, /  \\, f_b $, never exceeds 25 - 30\\%. Furthermore, $\\eta_\\ast$ peaks for Milky Way-sized galaxies ($M_h \\approx 10^{12} M_\\odot$), and declines rapidly on either side of the peak (e.g. Figure 2 in \\citealp{Guo2010}).  The second line of evidence comes from direct halo mass measurements, obtained through weak lensing or kinematics studies \\citep[e.g.][]{Dutton2010, Reyes2012}. For example, \\citet{Reyes2012} used stacked weak lensing measurements to estimate the average host halo mass of galaxies in different stellar mass bins, and found that $\\eta_\\ast$ never exceeds $\\approx 30$\\%. Direct halo mass measurements can circumvent a number of assumptions inherent in the application of abundance matching, but such techniques can presently only be applied to a restricted range of stellar mass ($M_\\ast \\approx 10^9 - 10^{11} \\: M_\\odot$), and are affected by their own set of systematics.  Stellar mass is not always the dominant baryonic component in a galaxy. In fact, the HI--to--stellar mass ratio (``HI fraction''; $f_{HI} = M_{HI} / M_\\ast$) tends to increase with decreasing stellar mass, and HI often dominates the baryonic content of low-mass galaxies. The transition from stellar-mass--dominated to HI--dominated systems takes place at $M_\\ast \\approx 10^{10} M_\\odot$ for HI-selected samples \\citep[e.g.][see also Fig. \\ref{fig:gsmr} in this work]{Huang2012b}, or at $M_\\ast \\lesssim 10^{9.5} \\: M_\\odot$ for optically-selected samples \\citep[e.g.][]{Catinella2010}. As a result, it is presently not clear what is the behavior of the ``baryon retention fraction'' $\\eta_b = (M_b/M_h) \\, / \\, f_b$ in low-mass galaxies, when both stars and cold gas are taken into account. In particular it is not well understood whether the very low average value of $\\eta_\\ast$ inferred for low-mass halos is a result of poor retention of baryonic material, of the low efficiency of gas-to-stars conversion, or of a combination of both.  For example, \\citet{Baldry2008} argue that the increasing gas fraction in low-mass galaxies should approximately offset the decreasing stellar-to-halo mass ratio, and result in a roughly constant $\\eta_b \\approx 10$\\%. This conclusion was based on an indirect estimate of the cold gas content of galaxies, based on the average $f_{HI} - M_\\ast$ relation observed in a set of samples in the literature. An early work by \\citet{SP1999}, based on the same indirect method, also reached a qualitatively similar conclusion. \\citet{Evoli2011} found an approximately constant $\\eta_b$ at the low-mass end using a different indirect method, which involves the comparison of the stellar and HI mass distributions of two different galaxy samples. These results would imply that low-mass galaxies are relatively efficient at retaining baryonic mass, but very inefficient in converting their  gas into stars. This conclusion, however, would require a ``steep'' HI mass function (HIMF) in the local universe, in contrast to what is measured \\citep{Zwaan2005, Martin2010}. Moreover, the recent work of \\citet{Rodriguez2011}, also based on using the average $f_{HI} - M_\\ast$ relations for blue and red galaxies separately, found no signs for a flat $\\eta_b$ at low masses.  In this article we \\textit{directly} measure the abundance of galaxies as a function of their ``baryonic mass'' (throughout this article the term baryonic refers to the combined stellar and atomic gas components of galaxies, and baryonic mass is calculated as $M_b = M_\\ast + 1.4 \\, M_{HI}$, where the 1.4 factor accounts for the presence of helium). We use optical data from the seventh data release of the SDSS (SDSS DR7) to estimate stellar masses, and HI-line flux measurements from the Arecibo Legacy Fast ALFA\\footnotemark{} (ALFALFA) survey to measure atomic gas masses. The resulting distribution, referred to hereafter as the baryon mass function (BMF) of galaxies, can be used in abundance matching to derive a robust $\\eta_b$ - $M_h$ relation. In order to investigate sample selection effects, we employ both an HI-selected and an optically-selected sample drawn from the \\textit{same} volume to derive the mass distributions for the stellar, atomic hydrogen and baryonic components.   \\footnotetext{The Arecibo L-band Feed Array (ALFA) is a 7-feed receiver operating in the L-band ($\\approx 1420$ MHz), installed at the Arecibo Observatory.}  The paper is organized as follows: in section \\ref{sec:data}, we introduce the datasets used to measure the stellar, HI and baryon mass distributions. We describe the methodology used to measure atomic hydrogen masses and we estimate stellar masses for our galaxy samples. In section \\ref{sec:bmf}, we present our measurements of the SMF, HIMF \\& BMF from both the HI-selected and the optically-selected samples, and compare them against one another as well as against other published results. In section \\ref{sec:biases} we consider the impact of possible systematics on our measurements, such as stellar mass estimation method, distance uncertainties and the exclusion of some baryonic components (e.g. molecular gas) in the calculation of the BMF. In section \\ref{sec:bar_frac}, we present the $\\eta_\\ast$ - $M_h$ and $\\eta_b$ - $M_h$ relation in a $\\Lambda$CDM universe. In section \\ref{sec:conclusion}, we discuss the implications of the result and summarize our main conclusions. Throughout this paper, we use a Hubble constant of $H_0 = 70 \\: h_{70} \\;\\; \\mathrm{ km}\\, \\mathrm{s}^{-1}  \\mathrm{Mpc}^{-1}$. ",
        "conclusions": "}  We use optical data from the seventh data release of the Sloan Digital Sky Survey (SDSS DR7) and 21cm emission-line data from the Arecibo Legacy Fast ALFA (ALFALFA) survey to measure the ``baryonic mass'' (defined as $M_b = M_\\ast + 1.4 \\, M_{HI}$) of galaxies in the local universe, and determine the $z=0$ baryon mass function (BMF). We use both an HI-selected and an optically-selected sample (7195 and 22587 galaxies respectively) drawn from the same volume, in order to address the effects of sample selection on the derived distributions. We find that the main difference consists of the optically-selected stellar mass function (SMF) being systematically larger at high-masses than the HI-selected SMF, and find that this difference carries over to the high-mass end of the BMF (see Fig. \\ref{fig:opt_vs_hi_gsmf} \\& \\ref{fig:opt_vs_hi_bmf}).    We combine the obtained mass distributions with the halo mass function in a WMAP3 $\\Lambda$CDM cosmology, to obtain average values of $M_\\ast / M_h$ and $M_b / M_h$ as a function of halo mass (Fig. \\ref{fig:am1} \\& \\ref{fig:am2}). Our most important result is that low-mass halos seem to have very low galactic baryon fractions compared to the cosmic value ($f_b = \\Omega_b / \\Omega_m \\approx 0.16$), even when their atomic gas content is taken into account; for example, the average baryon fraction of halos with $M_h = 10^{10.3} \\, M_\\odot$ is just 2\\% of the cosmic value ($\\eta_b \\approx 0.02$), and displays a monotonically decreasing trend. This result contrasts with previous indirect measurements of the BMF \\citep{Baldry2008, Evoli2011}, which pointed to an approximately constant value of $\\eta_b \\approx 0.10$ at the low halo-mass end.   Such very low values of $\\eta_b$ are difficult to reconcile with current models of galaxy formation. Photoionization heating in the early universe suppresses the baryonic content of halos only at $M_h \\lesssim 10^{10} \\, M_\\odot$ \\citep{Okamoto2008}, but this mass is more than an order of magnitude smaller than what is required by our result. Therefore, additional feedback mechanisms, such as baryon blowout by supernova explosions, must be present and must be extremely efficient. It is not yet clear whether hydrodynamic simulations or observational results can accommodate such intense galactic outflows in low mass halos. As a result, the observed $\\eta_b$ - $M_h$ relation remains difficult to explain, and may represent a challenge to our understanding of galaxy formation and/or the properties of dark matter."
    },
    "1208/1208.5645_arXiv.txt": {
        "abstract": "It is known that the diffuse H$\\alpha$ emission outside of bright \\ion{H}{2} regions not only are very extended, but also can occur in distinct patches or filaments far from H II regions, and the line ratios of {[}\\ion{S}{2}{]} $\\lambda$6716/H$\\alpha$ and {[}\\ion{N}{2}{]} $\\lambda$6583/H$\\alpha$ observed far from bright \\ion{H}{2} regions are generally higher than those in the \\ion{H}{2} regions. These observations have been regarded as evidence against the dust-scattering origin of the diffuse H$\\alpha$ emission (including other optical lines), and the effect of dust scattering has been neglected in studies on the diffuse H$\\alpha$ emission. In this paper, we reexamine the arguments against dust scattering and find that the dust-scattering origin of the diffuse H$\\alpha$ emission cannot be ruled out. As opposed to the previous contention, the expected dust-scattered H$\\alpha$ halos surrounding \\ion{H}{2} regions are, in fact, in good agreement with the observed H$\\alpha$ morphology. We calculate an extensive set of photoionization models by varying elemental abundances, ionizing stellar types, and clumpiness of the interstellar medium (ISM) and find that the observed line ratios of {[}\\ion{S}{2}{]}/H$\\alpha$, {[}\\ion{N}{2}{]}/H$\\alpha$, and \\ion{He}{1} $\\lambda$5876/H$\\alpha$ in the diffuse ISM accord well with the dust-scattered halos around \\ion{H}{2} regions, which are photoionized by late O- and/or early B-type stars. We also demonstrate that the H$\\alpha$ absorption feature in the underlying continuum from the dust-scattered starlight (``diffuse galactic light'') and unresolved stars is able to substantially increase the {[}\\ion{S}{2}{]}/H$\\alpha$ and {[}\\ion{N}{2}{]}/H$\\alpha$ line ratios in the diffuse ISM. ",
        "introduction": "The diffuse H$\\alpha$ emission outside of bright \\ion{H}{2} regions is ubiquitous in late-type galaxies and is generally believed to probe the warm ionized medium (WIM; also called diffuse ionized gas, DIG), which may be a major component of the interstellar medium (ISM) of our Galaxy and other late-type galaxies \\citep{Reynolds91,Walterbos96,Dettmar2000,Rand2000,Hidalgo-Gamez2005,Haffner2009}. It has been argued that the WIM is mainly photoionized by ionizing radiation (Lyman continuum; Lyc) that leaks out of bright \\ion{H}{2} regions associated with O stars \\citep{Mathis1986,Mathis2000,Domgorgen94,Sembach2000,Wood2004,Wood2010}, although it is not clear how the Lyc can penetrate the diffuse \\ion{H}{1} that is observed everywhere in the galaxies \\citep[e.g., ][]{Seon2009}.  Previous studies of the diffuse H$\\alpha$ and other optical emission lines concluded that the effect of dust scattering is negligible \\citep{Reynolds1985a,Walterbos1994,Ferguson96a,Ferguson96b,Wood1999}. However, it was recently revealed that the H$\\alpha$ excess intensity in a number of high-latitude clouds is due to scattering by interstellar dust of H$\\alpha$ photons originating elsewhere in the Galaxy \\citep{Mattila2007,Lehtinen2010,Witt2010}. More recently, \\citet{Seon2011a,Seon2011b} presented a strong connection between the diffuse H$\\alpha$ and far-ultraviolet (FUV) continuum backgrounds, which likely suggests a common radiative transfer mechanism responsible for a substantial fraction of both H$\\alpha$ line and FUV continuum photons. Through a comparison of the H$\\alpha$ and FUV continuum backgrounds, we proposed that the diffuse H$\\alpha$ emission at high latitudes may be dominated by late O- and/or early B-type stars (hereafter, late OB stars) and dust scattering \\citep{Seon2011b}. Furthermore, \\citet{Brandt2012} detected emission lines H$\\alpha$, H$\\beta$, {[}\\ion{N}{2}{]} $\\lambda$6583, and {[}\\ion{S}{2}{]} $\\lambda$6716 at high Galactic latitudes, whose intensities are correlated with the 100 $\\mu$m intensity, providing an independent confirmation of the importance of dust-scattered components in the optical emission lines. If the dust scattering indeed significantly contributes to the diffuse H$\\alpha$ emission and the effect of dust scattering is not properly taken into account, the energy requirement for the ionization and volume filling fraction of the WIM will be overestimated. Also, the widely-applied practice of using the all-sky map of diffuse H$\\alpha$ intensity \\citep{Finkbeiner2003} as a template for both the intensity and small-scale structure of the Galactic foreground free-free emission in cosmic microwave studies \\citep[e.g., ][]{Dickinson2009,Gold2011} needs to be reevaluated, if dust scattering contributes significantly to the intensity and spatial structure of diffuse H$\\alpha$ light outside the actual \\ion{H}{2} regions.  The recent findings about the importance of dust scattering motivated us to reexamine the rationales behind previous conclusions against the dust-scattering origin of the diffuse H$\\alpha$ emission. It was also necessary to confirm our previous results that nebulae photoionized by late OB stars are capable of reproducing the typical values of the {[}\\ion{N}{2}{]} $\\lambda$6583/H$\\alpha$ and {[}\\ion{S}{2}{]} $\\lambda$6716/H$\\alpha$ line ratios.  The rejection of dust scattering as a significant contributor to the diffuse H$\\alpha$ and other optical line emissions is mostly based on two arguments. The first argument is based on the H$\\alpha$ morphology in the surroundings of \\ion{H}{2} regions. \\citet{Ferguson96a,Ferguson96b} assumed that in \\emph{R}-band continuum images, bright OB associations would show \\emph{R}-band continuum halos as prominent as the H$\\alpha$ halos around \\ion{H}{2} regions, if the H$\\alpha$ halos around \\ion{H}{2} regions are caused by dust scattering. However, they found no corresponding halos around bright OB associations in the \\emph{R}-band images. Additionally, light scattered by dust was assumed to be usually concentrated around \\ion{H}{2} regions, but the diffuse H$\\alpha$ emission is often observed in distinct patches or filaments far from \\ion{H}{2} regions \\citep[e.g., ][]{Hoopes1996}. Consequently, the extended H$\\alpha$ emission outside of bright \\ion{H}{2} regions was regarded as evidence of leakage of Lyc photons rather than as the dust scattering of H$\\alpha$ photons. Second, the optical line ratios of {[}\\ion{N}{2}{]}/H$\\alpha$ and {[}\\ion{S}{2}{]}/H$\\alpha$ observed outside of \\ion{H}{2} regions have been found to be generally higher than the corresponding ratios in bright \\ion{H}{2} regions \\citep{Reynolds1985a,Reynolds1987,Walterbos1994}. It was thereby concluded that the dust-scattering effect could be ignored in spite of the fact that bright \\ion{H}{2} regions are usually surrounded by a large amount of molecular gas and dust.  In the present study, by using radiative transfer models, we reexamine whether these arguments are indeed valid in rejecting the significance of dust scattering in the diffuse optical emission. First, we argue that the assertion based on the halo extent or morphology is clearly wrong. Instead, we show that the global H$\\alpha$ morphology in face-on galaxies can be readily explained by dust scattering. Second, we find that the observed line ratios of {[}\\ion{N}{2}{]}/H$\\alpha$ and {[}\\ion{S}{2}{]}/H$\\alpha$ outside of \\ion{H}{2} regions can be explained reasonably well by the dust-scattered halos around the \\ion{H}{2} regions photoionized by late OB stars, which is consistent with the main conclusion of \\citet{Seon2011b}. We demonstrate these results by calculating photoionization models of the radiation-bounded \\ion{H}{2} regions and putting the resulting emissivity distributions of the optical lines into the radiative transfer models as input sources for a dust scattering simulation, thereby showing that the dust-scattering origin of the diffuse H$\\alpha$ emission cannot be ruled out. Lastly, we also point out that the H$\\alpha$ absorption line in the dust-scattered starlight (so called ``diffuse galactic light,'' DGL) and the stellar continuum of unresolved stars can significantly increase the observed values of the line ratios.  This paper is organized as follows. In Section 2, we discuss the morphology and the surface brightness distribution of H$\\alpha$ halos around \\ion{H}{2} regions caused by dust scattering. Section 3 investigates one-dimensional and three-dimensional photoionization models, and the line ratios resulting from the models. In Section 4, we discuss the effect of the underlying Balmer absorption line on the optical line ratios. Additional evidence supporting the results are described in Section 5. Section 6 summarizes the results. ",
        "conclusions": "In this section, we present additional support for our results and argue the necessity of reconsidering the dust-scattering origin of the diffuse H$\\alpha$ emission more seriously. In fact, we should note that there is no direct evidence, given the results presented in the previous sections, supporting the contention that most of the diffuse H$\\alpha$ emission originates from in situ ionized gas. The previous studies of \\citet{Mattila2007}, \\citet{Witt2010}, and \\citet{Seon2011b} have mainly dealt with the dust-scattered H$\\alpha$ photons at high Galactic latitudes. The present results strongly suggest the significance of dust scattering in the diffuse H$\\alpha$ emission not only at high-latitude clouds but also near \\ion{H}{2} regions.   \\subsection{Dust-Scattered H$\\alpha$ Halos}  There is observational evidence that at least some parts of the diffuse H$\\alpha$ emission outside of classical bright \\ion{H}{2} regions may be in fact caused by dust scattering. If the dust scattering is the dominant source of the diffuse H$\\alpha$ halos around the bright \\ion{H}{2} regions, the halos would not be observable or would be very faint in a radio recombination continuum that is almost free from dust-scattering. In this way, at least some part of the extended H$\\alpha$ halo of the Orion \\ion{H}{2} region was attributed to the dust-scattered halo \\citep{Subrahmanyan2001,ODell2009}. More direct evidence of dust-scattered origin of the H$\\alpha$ halo can be provided by observations of the H$\\alpha$ polarization. \\citet{Topasna1999} attempted to detect H$\\alpha$ polarizations by selective extinction in the Monoceros supernova remnant, the Rosette Nebula, and the North America Nebula. While these were not detected, H$\\alpha$ polarizations by scattering in dust shells around the Rosette Nebula and the North America Nebula were instead observed.  Additional support for the importance of dust scattering in the diffuse H$\\alpha$ emission is provided from the large scale surveys of our Galaxy. \\citet{Kutyrev2001,Kutyrev2004} carried out a pilot survey of the Galactic plane with 2.17 $\\mu$m Br$\\gamma$ hydrogen RL, which is relatively free from dust-scattering, and found a typical filling factor of a Br$\\gamma$-emitting gas of $\\sim1$ \\%, which is much smaller than the value derived for H$\\alpha$. Recent observations using the Wilkinson Microwave Anisotropy Probe have found that the ratio of the free\\textendash{}free radio continuum to H\\textgreek{a} is surprisingly low in the WIM \\citep{Davies2006,Dobler2008a,Dobler2008b,Dobler2009,Gold2011}. These discrepancies likely suggest the significant role of dust scattering in producing the diffuse H$\\alpha$ emission outside of bright \\ion{H}{2} regions. However, we note that this discrepancy is not just present in regions where dust scattering may contribute, but even in regions that appear to be rather dust free, e.g., the Gum nebula. The cause of this observed fact may therefore be more complex and involve other factors besides scattering.  It is also of interest to note that a correlation between H$\\alpha$ and the anomalous dust emission at $\\sim30$ GHz, which is attributed to spinning very small dust grains, is found in the WMAP data \\citep{Dobler2008a,Dobler2008b}. The spectrum of H$\\alpha$-correlated microwave emission does not follow the expected free-free spectrum from the WIM, but the electric dipole emission from rapidly rotating dust grains. Since then, there have been a number of studies showing the association of peaks in the anomalous dust emission with individual interstellar clouds \\citep[e.g., ][]{Vidal2011}. The spinning dust microwave emission is simply the signal from the smallest component of the grain size distribution in interstellar space. The correlation of the anomalous dust emission with H$\\alpha$ is explained most readily, if a significant fraction of the H$\\alpha$ is also coming from dust via scattering.  The photoionization models of the diffuse H$\\alpha$ emission assume no significant absorption of the Lyc \\citep{Domgorgen94}. A widely accepted view for the almost-free escape of Lyc from the bright \\ion{H}{2} regions is to assume a highly clumpy or turbulent density structure of ISM with a lot of almost-empty space. \\citet{Ferguson96a,Ferguson96b} found a clear correlation between the diffuse H$\\alpha$ intensity and the surface brightness of \\ion{H}{2} regions on both large and small scales. This correlation was regarded to be evidence that the Lyc leakage is the dominant source responsible for producing the diffuse H$\\alpha$ photons. However, as we demonstrated, the radiative transfer models for the dust-scattered halos around single \\ion{H}{2} regions can reproduce the basic properties of the H$\\alpha$ halos surrounding \\ion{H}{2} regions in face-on galaxies. An advantage of the dust-scattering scenario investigated in the present study over the in situ ionization scenario would be the fact that the diffuse H$\\alpha$ fraction of $\\approx50$\\% can be explained naturally.  Since the analysis was based on the halos surrounding single \\ion{H}{2} regions, it is necessary to investigate if dust scattering can indeed explain the global H$\\alpha$ morphology, which may be an overlapping of the dust-scattered halos surrounding individual \\ion{H}{2} regions, of the face-on galaxies. In this regard, we point out that the analyses of \\citet{Zurita02} and \\citet{Seon2009}, which were initially intended to explain the global H$\\alpha$ morphology in terms of photoionization models, can be interpreted as accounting for the morphology by the overlapped dust-scattered halos. \\citet{Zurita02} attempted to model the diffuse H$\\alpha$ morphology seen in face-on galaxies assuming that the Lyc leakage from bright \\ion{H}{2} regions would be the predominant source of the diffuse H$\\alpha$. They assumed that a constant fraction of the Lyc leaks from bright \\ion{H}{2} regions and transferred through the two dimensional ISM disk with no significant absorption, and claimed that the diffuse H$\\alpha$ surface brightness of NGC 157 can be reproduced reasonably well with this scenario. \\citet{Seon2009} extended the analysis to three dimensions assuming an exponential disk, and performed a similar analysis for a face-on galaxy, M 51. We then concluded that the effective hydrogen density required to explain the diffuse H$\\alpha$ emission with the ``standard'' scenario was too small ($\\sim10^{-5}$ of the generally known value) to be reconciled with the effective density inferred through the turbulence properties of ISM. In other words, unlikely extreme topology of the ISM is required to explain the diffuse H$\\alpha$ morphology by photoionization in a turbulent or clumpy medium.  These analyses, however, did not entail self-consistent photoionization modeling, as noted in \\citet{Wood2010}. The photoionization model should be iteratively performed until the ionization fraction and temperature at each points converge. Instead, the penetration of photons through ISM was characterized using an $\\exp(-\\tau)$ factor, after being diluted by a geometric factor $1/r^{2}$, where $\\tau$ is the effective optical depth, and the photons interacting at a point were assumed to be re-emitted isotropically. This radiative transfer process in fact describes a single dust-scattering with an arbitrary dust extinction cross-section $\\sigma_{{\\rm eff}}$ and an isotropic scattering phase function (asymmetry phase factor $g=0$). In our simulations, it was found that the H$\\alpha$ luminosity from the diffuse region (DIG) is approximately half of the total H$\\alpha$ luminosity originating from \\ion{H}{2} regions. In the context of a single dust-scattering, this is equivalent to assuming the dust albedo of $a\\sim0.5$. From the analysis, it was also found that the effective cross-section per hydrogen nucleus is given by $\\sigma_{{\\rm eff}}\\approx(2-4)\\times10^{-22}$ cm$^{2}$ for a reasonable hydrogen density. Note that the obtained effective cross-section is a factor of $\\approx10^{-5}$ smaller than the photoionization cross-section of Lyc, but accords well with the dust-extinction cross-section of H$\\alpha$ photons of $\\sigma_{{\\rm ext}}=3.8\\times10^{-22}$ cm$^{2}$/H \\citep{Draine03}. The absorption coefficient obtained by \\citet{Zurita02} is also consistent with the value. Therefore, the results indicate that the diffuse H$\\alpha$ morphologies in the face-on galaxies accord well with the dust-scattered halos of H$\\alpha$ photons, contrary to the claim of \\citet{Zurita02} that the analysis supports the photoionization by Lyc leaked out of bright \\ion{H}{2} regions.  In Section 2, we found that the surface brightness of the individual dust-scattered halos can be approximated well by the formula $I(r)\\propto\\exp(-\\tau)/r^{2}$, which is in fact the very expression used in \\citet{Zurita02} and \\citet{Seon2009}. These results seem to strongly suggest the dust-scattering origin of the diffuse H$\\alpha$ emission in the face-on galaxies. Obviously, the radial profile is not consistent with those of the radiation-bounded \\ion{H}{2} regions. However, it is not clear at this stage whether or not the radial profile of the photoionized H$\\alpha$ emission line that is predicted from the density-bounded models accords with $I(r)\\propto\\exp(-\\tau)/r^{2}$. We plan to examine the radial profiles of density-bounded \\ion{H}{2} regions in the future. We also need to develop more realistic models of the global H$\\alpha$ morphologies of face-on galaxies in the contexts of not only dust scattering, but also photoionization.   \\subsection{Optical Line Ratios}  The {[}\\ion{N}{2}{]}/H$\\alpha$ and {[}\\ion{S}{2}{]}/H$\\alpha$ line ratios are observed to be higher in the diffuse H$\\alpha$ regions than in bright \\ion{H}{2} regions \\citep{Haffner2009}. This was believed to be the strongest evidence against the dust-scattering origin of the diffuse H$\\alpha$ emission, because the dust scattering would keep the line ratios constant. The enhanced optical line ratios outside of \\ion{H}{2} regions have been generally explained by a dilute radiation field transferred from O stars within \\ion{H}{2} regions to the diffuse ISM \\citep{Mathis1986,Sokolowski91,Domgorgen94}.  However, we note that the argument to reject dust scattering has been based solely on comparisons with the line ratios from bright \\ion{H}{2} regions. If the line ratios from the ionized nebulae due to the late O- or early B-type stars are compared with the typical line ratios in the diffuse H$\\alpha$ regions, the present scenario seems to explain the higher line ratios than those in bright \\ion{H}{2} regions equally as well as the previous photoionization models do. It was also found that the {[}\\ion{S}{2}{]}/H$\\alpha$ ratios observed in \\ion{H}{2} regions around late-type stars tend to be higher than those found around early-type stars (Figure 3 of \\citealp{Reynolds88}). \\citet{ODell2011} pointed out a striking similarity of the line ratios between Barnard\\textquoteright{}s Loop, the Orion\\textendash{}Eridanus Bubble, and the typical WIM samples, which can be explained by the photoionized nebulae due to stars with temperature of $\\lesssim35,000$ K corresponding to late O-types (see also, \\citealt{Seon2011b}). This also supports the present conclusion.  Weak {[}\\ion{O}{3}{]} $\\lambda$5007/H$\\alpha$ and \\ion{He}{1}/H$\\alpha$ emission line ratios in the galactic disks also indicate that the spectrum of the diffuse interstellar radiation field is significantly softer than that from the average Galactic O star population \\citep{Reynolds95,Madsen06}. These line ratios provide the strongest constraint on the spectral type of the original ionizing source of the diffuse H$\\alpha$ and other optical emission lines. However, {[}\\ion{O}{3}{]} $\\lambda$5007/H$\\beta$ or {[}\\ion{O}{3}{]} $\\lambda$5007/H$\\alpha$ in several halos of edge-on galaxies are found to increase with height above the galactic plane \\citep{TuellmannDettmar2000,MillerVeilleux2003}. The rise may indicate that other mechanisms are needed.  \\citet{Seon2011b} estimated how a large portion of the ionizing power required for the H$\\alpha$ background in our Galaxy can be provided by late OB stars, and concluded that O9 and later-type stars can account for at least one-half of the required recombination rate of $r_{{\\rm G}}\\approx2.5\\times10^{6}$ s$^{-1}$ cm$^{-2}$. We note that the earliest stellar type consistent with the observed \\ion{He}{1}/H$\\alpha$ line ratio is O8 from our analyses, and the fraction of Lyc luminosities of O8 and O9 stars in the total Lyc luminosity due to O stars is $\\sim18$\\% (Table 1 in \\citealp{Terzian1974}). Including O8 stars, late OB stars can provide the recombination rate of $(2.3-6.2)\\times10^{6}$ s$^{-1}$ cm$^{-2}$. It is therefore clear that most of the required ionizing power in the Galaxy would be explained with late OB stars. In fact, there is a relative lack of early O stars in the solar neighborhood, as noted in \\citet{Brandt2012}. Therefore, the diffuse optical emission lines in our Galaxy are consistent with those produced by late OB stars and their scattered light from the local ISM.  In our models using the MOCASSIN code, the typical values of the {[}\\ion{N}{2}{]}/H$\\alpha$ and {[}\\ion{S}{2}{]}/H$\\alpha$ ratios were well reproduced with most of the abundances except the ISM abundances. With the CLOUDY code, abundances close to those of WNM were required. \\citet{Sembach2000} and \\citet{Mathis2000} also reproduced the observed line ratios assuming the WNM abundances. In our previous paper \\citep{Seon2011b}, we mistakenly stated that the elemental abundances required to explain the line ratios in the WIM are close to the B star abundances both in \\citet{Sembach2000} and in \\citet{ODell2011}. This should be read as the WNM abundances instead of the B star abundances.  The NLTE stellar models may produce significantly different amount of Lyc photons from the LTE models. Nevertheless, we obtained consistent results from two different model atmospheres (Kurucz or WMBASIC) in Section 3.2. Note that it is the combination of the line blanketing and spherical geometry, not exclusively the NTLE aspect, that is crucial in determining the spectral energy distribution (SED) and Lyc luminosity of stellar models \\citep{Aufdenberg1998}. \\citet{Aufdenberg1999} also found that a significant difference between the spherical and plane-parallel models is found only at surface gravities below $\\log g_{*}=3.5$, corresponding to giant stars. We also note that the calibration scale of stellar parameters (temperature and gravity) as a function of spectral type varies when different model atmosphere codes are used \\citep{Martins2005,Simon-Diaz2008}. \\citet{Simon-Diaz2008} investigated the impact of the modern model atmospheres (including plane-parallel stellar models) of dwarf stars on \\ion{H}{2} regions. They concluded that the predicted SEDs in the energy range of $13-30$ eV, which are most important in producing H$\\alpha$, {[}\\ion{N}{2}{]}, {[}\\ion{S}{2}{]}, and \\ion{He}{1} lines, are in good agreement between the codes when the models are compared using the calibration scales appropriate to each code. Therefore, our main results obtained by using dwarf stars are not sensitive to the adopted atmosphere models. Since the number of giant stars is relatively smaller than that of dwarfs, it is also unlikely that the main conclusions will be significantly altered, even when giants are taken into account. Fully three-dimensional simulations with a set of NLTE input spectra, including the effect of giant stars, may be needed for a more detailed study.  The results of \\citet{Topasna1999} are worth noting in attempting to understand the line ratios of the diffuse H$\\alpha$ regions. The {[}\\ion{S}{2}{]}/H$\\alpha$ ratio observed in the Rosette Nebula increased from $\\lesssim0.2$ to $\\sim0.5$ with radius out to the boundary of the \\ion{H}{2} region at $r\\lesssim35'$, which was identified by polarization and 4850 MHz radio continuum emission. Beyond the \\ion{H}{2} region boundary ($r\\approx35'-60'$) where dust scattering is dominant, there were no large variations in the ratio as a dust-scattered halo would preserve the line ratio at the \\ion{H}{2} region boundary (see Fig.~4.27 of \\citealp{Topasna1999}). The $\\zeta$ Oph \\ion{H}{2} region also shows the same trend in the {[}\\ion{N}{2}{]}/H$\\alpha$ and {[}\\ion{S}{2}{]}/H$\\alpha$ ratios (Fig.~2 of \\citealp{Wood2005}). This trend is exactly what is expected from the dust-scattered halo. However, farther out beyond these regions with constant ratios, the line ratios increased to higher values. These higher ratios may indicate the existence of other sources close to the nebulae to elevate the line ratios, as will be discussed in the following paragraph.  It is thought that the line ratios increase as the H$\\alpha$ intensity decreases. This trend does not appear to be consistent with the dust scattering scenario, because dust scattering should preserve the ratios. We therefore need to explain the trend in the context of the dust scattering scenario. Note that early O stars in association are more concentrated at the galactic plane. Near an early O star, the dust-scattered H$\\alpha$ halo would be dominated by the \\ion{H}{2} region due to the O star and the H$\\alpha$ intensity will be very high. The {[}\\ion{N}{2}{]}/H$\\alpha$ and {[}\\ion{S}{2}{]}/H$\\alpha$ ratios would be then relatively low in the region. At distances far from the O star, the H$\\alpha$ intensity will be dominated by the dust-scattered light originating from the \\ion{H}{2} regions surrounding late OB stars, which are more abundant than early O stars, and the line ratios will then be increased. As we demonstrated using the clumpy models, some part of the enhancement or fluctuation in the line ratios may be attributed to the variation of clumpiness from sightline to sightline. Note that the {[}\\ion{S}{2}{]}/H$\\alpha$ ratio of $\\sim0.23$ in the $\\zeta$ Oph \\ion{H}{2} region \\citep{Wood2005} is close the ratio expected from the C1 model, while the {[}\\ion{S}{2}{]}/H$\\alpha$ ratio of $\\sim0.5$ in the Rosette Nebula \\citep{Topasna1999} needs a clumpier model than the C2 case. In LDN\\,1780, the ratio of $\\sim0.16$ \\citep{Witt2010} is somewhere in the range between the U and C1 model. However, the highest line ratios may need other explanations. The highest line ratios can result from non-ionizing heating sources, as discussed in \\citet{Seon2011b}, and/or from the underlying Balmer absorption line in the stellar continuum of relatively late-type stars, as discussed in Section 4.  The heating sources to elevate the line ratios may include shocks, photoelectric heating, turbulent mixing layers, and/or galactic fountain gas \\citet{Reynolds_Cox1992,Slavin1993,Collins_Rand2001,Raymond1992,Shapiro1993}. Recently, \\citet{Flores-Fajardo2011} proposed hot low-mass evolved stars (HOLMES), which may be plentiful in the thick disks and lower halos of galaxies, to explain the observed increase of {[}\\ion{O}{3}{]}/H$\\beta$, {[}\\ion{O}{2}{]}/H$\\beta$ and {[}\\ion{N}{2}{]}/H$\\alpha$ with increasing distance to the galactic plane in an edge-on galaxy, NGC\\,891. We also note that the \\ion{H}{2} region associated with the B0.5+sdO binary $\\phi$ Per is elevated in {[}\\ion{S}{2}{]}/H$\\alpha$ \\citep{Madsen06}. These stars produce a much harder radiation field than massive OB stars, and are therefore able to produce \\ion{H}{2} regions with high electron temperatures. It is possible that all these supplementary sources may play some roles in enhancing the line ratios and their contributions might significantly vary from sightline to sightline. The optical emission lines due to these sources providing the highest line ratios should also be scattered into more extended regions than those originally produced by the sources. It should be noted that the conventional photoionization models by O stars in the galactic plane also need these heating sources to explain the highest ratios \\citep{Reynolds1999}. Therefore, the present dust scattering scenario can be an alternative explanation for the diffuse optical H$\\alpha$ emission that is equally as plausible as the in situ ionization scenario.  In Section 4, we argued that the {[}\\ion{N}{2}{]}/H$\\alpha$ line ratio (probably, {[}\\ion{S}{2}{]}/H$\\alpha$ as well) in the diffuse ISM regions seems to be consistent with that in bright \\ion{H}{2} regions, if the underlying Balmer absorption lines are taken into account. It was also demonstrated that the line ratios should indeed increase at the faint H$\\alpha$ regions outside bright \\ion{H}{2} regions, if the diffuse H$\\alpha$ emission is predominantly caused by dust scattering. We note that a similar argument to explain the anti-correlation of the line ratios with H$\\alpha$ intensity can be applied even to the in situ photoionization scenario. Then, the non-ionizing heating sources proposed to explain the highest line ratios may not be required anymore, regardless of what would be the main origin of the diffuse H$\\alpha$ emission, dust-scattered light or in-situ ionized gas. However, the dust scattering origin is more preferable if this is indeed the case."
    },
    "1208/1208.3643_arXiv.txt": {
        "abstract": "We present VLBI and archival Karl G. Jansky Very Large Array (VLA) and Westerbork Synthesis Radio Telescope (WSRT) observations of the radio afterglow from the gamma-ray burst (GRB) of 2003 March 29 (GRB 030329) taken between 672 and 2032 days after the burst.  The EVLA and WSRT data suggest a simple power law decay in the flux at 5 GHz, with no clear signature of any rebrightening from the counter jet. We report an unresolved source at day 2032 of size $1.18\\pm0.13$ mas, which we use in conjunction with the expansion rate of the burst to argue for the presence of a uniform, ISM-like circumburst medium.   A limit of $< 0.067$ mas yr$^{-1}$ is placed on the proper motion, supporting the standard afterglow model for gamma-ray bursts. ",
        "introduction": "Gamma-ray bursts (GRBs) are sudden, short-lived ($\\lesssim 10^2$ s) releases of energy of on the order $10^{51}$ erg \\citep{frailEA01, bloomEA03} in a region of $\\lesssim 100$ km.  The central engine, thought to be a black hole with an accretion disk formed either by the collapse of a massive evolved star (collapsar) for long GRBs ($\\gtrsim 2$s) or by the merger of two stellar remnants (black holes or neutron stars) for short GRBs ($\\lesssim 2$s).  The large, sudden energy release produces an $e^+e^-$,$\\gamma$ fireball in the form of collimated jets \\citep{paczynski93}.  Variability in the rate and velocity at which material is ejected from the central engine leads to internal shocks within the jet when fast-moving shells catch up to slower ones.  The resulting $\\gamma$-ray emission is responsible for the so-called prompt emission \\citep{reesMeszaros94}.  Whereas prompt emission lasts for $\\lesssim$ a few minutes, GRB afterglow emission can last far longer, up to years in the radio part of the spectrum, e.g. \\citep{frailEA00, bergerEA03c, taylorEA04, pihlstromEA07}.  The long-lived afterglow emission is primarily synchrotron radiation is usually attributed to an interaction with the shocked external medium that exists behind the external forward shock \\citep{meszarosRees97}. More recently, perplexing features observed by Swift in the early X-ray afterglow have lent weight to a possible alternative scenario wherein a long-lived reverse shock decelerates slow ejecta at the back of the original outflow as it gradually catches up with the shocked external medium \\citep{genetEA07, uhmBeloborodov07}. Whether the afterglow arises from the forward external shock or from a long lived reverse external shock, an external shock origin is strongly supported by measurements of the radio afterglow image size and its temporal evolution for the late time radio afterglow of GRB 030329 \\citep{taylorEA04, orenEA04, granotEA05, pihlstromEA07}.  To explain the intensity and duration of observed GRB afterglow emission, the presence of a circumburst medium must be taken into account.  External shocks naturally arise when jet material interacts with the circumburst medium, converting kinetic energy of the fireball into particle energy and luminosity \\citep{meszarosRees93}.  These interactions are expected to be collisionless, mediated instead by the tangled and compressed magnetic fields that exist at the shock boundary.  Electrons accelerating along magnetic field lines at the shock boundary produce the power law emission spectrum characteristic of gamma-ray bursts \\citep{meszarosRees93, katz94}.  The spectral and temporal evolution \\citep{sariEA98, granotsari02} of GRB afterglows, then, are governed by such factors as the structure and dynamical evolution of the relativistic jet (for a review, see \\citealt{granot07}), but also by the environment (e.g., constant density versus a windlike density profile; \\citealt{chevalierLi00}).  A detailed study of the size evolution of a GRB afterglow would make it possible to better constrain the density profile of the circumburst medium.  Progenitors of long gamma ray bursts are expected to be associated with the collapse of evolved stars (collapsars) \\citep{woosley93, paczynski98, FryerEA99}.  At least some GRBs should therefore exist in gas-rich environments with isotropic windlike $\\rho(r) \\propto r^{-2}$ density profiles out to a few tenths of a parsec, where there is a region of roughly uniform density corresponding to the stellar wind reverse shock \\citep{ramirez-RuizEA05}.  Because the radio afterglow is detectable for months or years after the initial explosion, the jet would be expected to have sufficient time to traverse the stellar wind and enter the uniform shell of material that lies beyond.  Several long GRBs, however, have shown light curves that are more characteristic of a jet propagating through a uniform $\\rho(r) = \\rho_0$ medium for the entire observed duration of the radio afterglow.  More direct and less model-dependent methods of determining the radial density profile could be used to better constrain the progenitor as well as to test the validity of the standard fireball model.  In the simplest emission models, the flux density evolution of the GRB afterglow is related to the density of the medium in which the GRB is present \\citep{waxman97}.  The afterglow from GRB 030329 was monitored by \\citet{taylorEA04} and \\citet{pihlstromEA07} to obtain the first well-determined expansion rate for a GRB up to 247 and 806 days after the burst, respectively.  The apparent diameter of the burst afterglow was shown to increase from $0.065 \\pm 0.022$  mas to $0.172 \\pm 0.043$ mas between day 25 and 83.  Thereafter, the expansion rate decreased, with the burst reaching a size of $0.347 \\pm 0.090$ mas at $t = 806$ days.  The evolution of the mean apparent expansion speed of the afterglow image \\citep{pihlstromEA07} suggests a transition to non-relativistic expansion after about one year (see Fig.\\ \\ref{beta}). Models of the afterglow expansion predicted similar sizes at day 217 and day 806 in both the uniform density and the wind density profile case (see Fig.\\ \\ref{models}), making it necessary to observe the afterglow at later times so that the preferred model for the GRB 030329 circumburst environment could be unambiguously determined using this method.  Here we present observations of the GRB 030329 afterglow taken up to 5.5 years after the initial burst in order to shed light on the validity of the different afterglow models.  We also seek to place constraints on the density profile of the circumburst environment, the jet dynamics and structure, and the proper motion. ",
        "conclusions": "We present measurements of the 5 GHz flux density and image size of the GRB 030329 radio afterglow taken over a period of 5.5 years.  These observations clearly demonstrate that the expansion rate has decreased over time, with a transition to the non-relativistic regime at $\\sim$ 1 yr.  After approximately day 59, the afterglow flux density follows a power law of $F_\\nu \\propto t^{-\\alpha}$, with temporal index $\\alpha = 1.27\\pm 0.03$, which agrees with the value of $\\alpha = 1.23\\pm 0.03$ obtained by \\citep{pihlstromEA07}.  Using the method of \\citet{vanderHorstEA08}, the electron power law index $p$ is found to be $p = 2.24\\pm0.02$ for a uniform medium, which does not agree with their value of $p = 2.12\\pm0.02$.  The value determined for the electron power law index is found to be highly sensitive to the time range used to calculate it, making the temporal slope only capable of providing a quick estimate of the electron power law index.   A rebrightening of the source was expected as the counter jet became non-relativistic, however, no rebrightening was detected up to 5.5 years after the burst.  Numerical simulations suggest that any rebrightening in a windlike medium would be more difficult to detect as compared to a uniform medium.  However, a windlike medium is not favored by the measured evolution of the afterglow image size.  A possibility also exists for asymmetry between the jet and the counter jet.  If the counter jet had a lower initial energy or encountered a less dense external medium, then it would produce a smaller rebrightening than expected.  An upper limit of $0.067$ mas yr$^{-1}$ is found for the proper motion.  Consequently, the proper motion of the flux centroid is constrained to be smaller than the diameter of the image, which is consistent with the fireball model but not with the cannonball model.  The upper limit that we obtained for the source size at day 2032 (Fig.\\ \\ref{models}) is, by itself, insufficient to conclusively determine the nature of the circumburst medium through which the jet is propagating.  We argue that, on physical grounds, the mean expansion rate $\\langle\\beta_{\\rm app}\\rangle$ cannot increase between two epochs.  The value of $\\langle\\beta_{\\rm app}\\rangle$ at day 806 was $0.9\\pm0.2$, which corresponds to a maximum angular size at day 2032 of 0.90 mas, which is not consistent with a windlike medium.  We therefore argue that the expansion rate of the burst favors a uniform medium.  Moreover, a non-steady wind might produce an intermediate density profile with $k\\sim 1$ where $\\rho_{\\rm ext} = Ar^{-k}$. This might potentially enable us to reconcile between the model and the afterglow image size evolution that favor a more uniform external medium, and the radio light curve that favors a more wind-like external density because of the lack of a bump in the light curve due to the counter-jet. This will be tested in a follow-up paper.  Future work includes a comparison of our observations to detailed numerical simulations of the jet dynamics in different external density profiles \\citep{deColleEA12, deColleEA11}.  Additionally, direct measurements of the afterglow expansion rate will be conducted on any future gamma-ray burst that is sufficiently bright to be observed with VLBI."
    },
    "1208/1208.6299_arXiv.txt": {
        "abstract": "In this paper we describe the way the Astro-WISE information system (or simply Astro-WISE) supports the data from a wide range of instruments and combines multiple surveys and their catalogues. Astro-WISE allows ingesting of data from any optical instrument, survey or catalogue, processing of this data to create new catalogues and bringing in data from different surveys into a single catalogue, keeping all dependencies back to the original data. Full data lineage is kept on each step of compiling a new catalogue with an ability to add a new data source recursively. With these features, Astro-WISE allows not only combining and retrieving data from multiple surveys, but performing scientific data reduction and data mining down to the rawest data in the data processing chain within a single environment. ",
        "introduction": "\\label{s:intro}   The increasing challenge of data management in astronomy is created not only by the increasing volume of the data flowing in from the telescopes, but as well by the variety of new astronomical catalogs and surveys being created. Combined analysis of multi-survey datasets from different instruments has the potential to solve outstanding questions in a wide range of astronomical and astrophysical areas. These include the combining of optical and near-infrared surveys to identify objects of interest such as very high-redshift Quasars (SDSS-UKIDSS, see~\\cite{QSO} for example), unusual Brown Dwarfs~\\cite{BD} and ultra-compact binaries~\\cite{CVs}. The very same surveys can also be combined to constrain fundamental properties of our Universe: the nature of dark energy, the nature and distribution of dark matter via galaxy weak lensing, correlations in galaxy and QSO distributions and growth of large-scale structures. Complex relationships such as evolution of galaxies and their environment and nuclear activity require combined analysis from radio to X-rays (GAMA~\\cite{GAMA}, AEGIS~\\cite{AEGIS}, COSMOS~\\cite{COSMOS}, Coma Legacy Survey~\\cite{COMALS}, ATLAS3D~\\cite{ATLAS3D}, GOODS~\\cite{GOODS}, HUDF~\\cite{HUDF}). The challenge for any information system hosting these datasets is no longer archiving of data products, but providing abilities to perform data mining and data reprocessing on a massive scale.  The combining of a number of surveys and the data mining of the resulting combined survey is not a trivial task due to the volume of the data and to the particular requirements each user has for the combined catalogue. There are many ways to perform this task--from using the abilities provided by the survey itself to employing resources and software of Virtual Observatory.  The Virtual Observatory (VO) is the system of standards for publishing and accessing astronomical data developed by the International Virtual Observatory Alliance\\footnote{http://www.ivoa.net/}. At the same time VO is the collection of all data available according to these standards. They are published by a number of organizations, for example, the European Virtual Observatory community\\footnote{http://www.euro-vo.org/}. Recently VO standards were used not only to give data access but also to process the data itself by Canadian Advanced Network For Astronomical Research (CANFAR\\footnote{http://canfar.phys.uvic.ca}).  Despite the progress achieved via the VO, combining of astronomical data from different surveys remains a challenge. The problem with any of the methods is the limited volume of data which can be combined into a new catalogue or archive and the detached nature of the produced catalogue which is a separate entity without dependencies to the parent surveys. In addition, the user must arrange for the storage of the produced data and invent a way to automate the production of the new compiled catalogue.  These technical problems can been overcome, but not generally in the most efficient manner. Legacy research with large surveys often takes an ``\\textit{end-point}'' approach: the end products of the surveys (e.g., calibrated images and catalogues) are taken in by the end-user as starting point for analysis. Combining just the catalogue end-products of surveys with $\\leq 10^8$ entries is manageable using the infrastructure of a single user (``desktop science''). However, scientific requirements may lead to re-processing (e.g., going further backwards towards the raw data). This can include homogenisation in terms of photometry, astrometry, aperture/image quality. As soon as the the analysis involves (re)processing of data and, hence, returning to the original images, it requires both hardware and software infrastructure well beyond the desktop regime to deal with the avalanche of heterogeneous/complex data. This reprocessing of the survey data often requires the full survey operation data flow system including the quality control task. In turn, this means that the user has to possess the similar human and computing resources as the original data processing site.   The natural way to solve these problems is by allowing the user to access the original data processing infrastructure. This gives the user an ability to reprocess data partially or fully to create his own version of the survey's end product. Nevertheless, this will not solve the problem of reprocessing another survey's data which the user would like to join with the first one. This requires a data processing system ideally allowing processing of data regardless of the survey/telescope/instrument/filters. Such a system must be designed to be generic enough so that the data from many different instruments can be handled by it. This system must have a generic pipeline with an ability to port new pipelines for new instruments and surveys.   The Astro-WISE information system~\\cite{adass} is such system for optical/near-infrared wide-field imaging. The Astro-WISE approach to the porting of new surveys and instruments allows the user to create a combination of surveys and share it with other users. The key ingredient that allows the combination of different surveys and catalogues into new data products is a flexible common data model implemented in Astro-WISE. The ability to process the data inside Astro-WISE is defined by the level of integration of the external data in the Astro-WISE system. The deepest level of integration allows the processing and analysis of data from its most raw form, directly from the detector. At the shallowest level the scientific analysis starts in catalog-space using ingested external catalogues.  The difference between an astronomical data warehouse like the Virtual Observatory and Astro-WISE with its integration of wide range of surveys is the ability of (re-)processing and traceability of processing. The basis for this ability is a single, common data model specifically designed for this purpose. The integration takes place in an existing system that has a novel implementation of query language and tools for combining of catalogues. Astro-WISE provides the user with the necessary resources to do the job and allows the storage and sharing of the result of the compilation with both team members and the outside community.  In the next sections we describe the general design and components of the solution that allows Astro-WISE to be used as a platform to operate multiple surveys. The solution is based on the core principles of Astro-WISE: a common data model, data lineage and a modular approach to programming (see~\\cite{WISE}). ",
        "conclusions": "The Astro-WISE information system allows data handling for multi-survey operations and research within a single environment. The handling spans the from calibration of raw survey data to a wide range of post-calibration analysis (morphometry, photometric redshifts, variability etcetera). The special feature of the Astro-WISE information system is the ability to keep all processing dependencies inside the system. No matter how complicated the data lineage is, the user can understand how the final catalogue is compiled via the sequence of operations starting from the raw data, and can reprocess it using different processing parameters. Given that Astro-WISE keeps the quality parameters for each data item in the processing chain, it becomes possible to provide the user with everything necessary to accomplish multi-survey quality control and research inside a single system.  The system is scalable in terms of imaging data sources (and hence in data volume). The counter currently stands at 18 different instruments which are supported in Astro-WISE. The user can add new instruments and surveys, and reuse code to reduce the new data sets and combine them with existing survey products. The key method that allows this achievement is the description of wide-field imaging data and instruments using a common data model. Astro-WISE allows (re)processing and analysis of an unprecedented wide range of surveys and archives with all operations done on the multiple surveys in a uniform manner.  Another advantage of the common data model is that the same programs and interfaces can be reused on new data, reducing time the user should spent on an adaptation to the new instrument or survey. This also saves time for interface building for new data sets~\\cite{interface}.  Presently Astro-WISE accommodates all types of astronomical catalogues varying from radio to X-rays. Source extraction and subsequent modelling and analysis can be done on radio, near-infrared and optical imaging. Optical and near-infrared imaging data at raw or any process level can be (re)processed within the system. Next steps in the development of the common data model are adaptation to accomodate spectroscopic data to make it possible to combine optical and spectroscopic surveys. The Astro-WISE approach for information processing has proven to be generic enough to be applied for radio survey handling (LOFAR~\\cite{lofar}), optical multi-unit spectroscopy (Multi-Unit Spectroscopic Explorer) and beyond astronomy in the fields of medicine and artificial intelligence\\footnote{http://www.rug.nl/target}.  \\begin{acknowledgement} The authors are very pleased to acknowledge the pioneering work by Edwin A. Valentijn, Kor G. Begeman, Danny R. Boxhoorn, Erik R. Deul and Roeland Rengelink in creating the OmegaCAM data model. This work laid the basis for the common data model in Astro-WISE. \\end{acknowledgement}"
    },
    "1208/1208.4063_arXiv.txt": {
        "abstract": "Polar ring galaxies (PRGs) are peculiar systems where a gas-rich, nearly polar ring surrounds a host galaxy. They are the result of galaxy interactions that form mainly by tidal accretion of material from a gas rich donor galaxy. There is a number of formation mechanisms for PRGs: minor or major mergers, tidal accretion events, or direct cold gas accretion from filaments of the cosmic web. These objects can be used to probe the three-dimensional shape of dark matter haloes, provided that the ring is in equilibrium with the gravitational potential of the host galaxy. The polar ring galaxy, AM\\,2040-620, which has not yet been well studied, is the subject of this work. This galaxy contains an almost perpendicular warped ring and one possible companion galaxy to the NW. The radial velocity of this object is 3301$\\pm$65 Km s$^{-1}$ and is part of a group of fifteen possible polar ring galaxies, according to the literature. In order to better understand this system, images and long slit spectra were observed with the 1.60 m OPD/LNA telescope. In the I band image, the outer parts of the ring are not symmetrical. A disturbance in the Eastern side and a faint plume were detected. Two small satellites are located to the north. The bulge is elliptical but not perfectly symmetrical in this image. The B-band image shows material that extends beyond the ring in the western and eastern directions. After processing, the B-image shows that the possible companion galaxy 2MASX J20441668-6158092 has a tidally disturbed disk. Its radial velocity is unknown, but the spectroscopy, which is still under analysis, will furnish this information. ",
        "introduction": "Polar ring galaxies (PRGs) constitute a rare class of interacting systems and consist of an early-type, lenticular, elliptical or even spiral host galaxy, surrounded by a ring of gas, dust and young stars orbiting in a nearly polar plane (e.g. \\cite{1990AJ....100.1489W}, \\cite{1993A&A...268..103A}, \\cite{1997yCat..41290357F}, \\cite{2001MNRAS.322..689R}, \\cite{2006A&A...446..447R}).  In these objects, the velocities of the ring and the host galaxy are quite close, with an almost balanced state. The ring material appears to be in regular rotation about the galaxy center (e.g. \\cite{2004AJ....128.2013C}, \\cite{1987ApJ...314..457V}, \\cite{1997AJ....113..585A}) and is presumably stabilized in some way. According to Iodice (2002) \\cite{2002AJ....123..195I}, for a PRG to become stable, the formation of the ring must follow a certain logic, in which the ring size is strongly related to the amount of matter in the host galaxy (visible + dark).  The presence of two almost perpendicular angular momentum vectors cannot be explained through the collapse of a single protogalactic cloud; therefore a ``second event'' must have occurred in the formation history of these objects. However, no one knows for sure how these objects are formed. There are three formation mechanisms for PRGs \\cite{2003A&A...401..817B}: minor or major mergers, tidal accretion events or direct cold gas accretion from cosmic web filaments. These objects can be used to probe the three-dimensional shape of dark matter haloes, provided that the ring is in equilibrium with the gravitational potential of the host galaxy.  In this paper, we study the PRG AM\\,2040-620. The galaxy has an elliptical bulge. Its with radial velocity is 3301$\\pm$65 km s$^{-1}$ (this paper), very similar (within the errors) to the value measured from HI data by van Driel (2002) \\cite{2002A&A...386..140V} (3335$\\pm$24 km s$^{-1}$). The ring is not radially thick and presents a warped appearence.   \\begin{figure} \\centering \\includegraphics[width=5cm]{fenda} \\caption{B-band image of NGC\\,5122 with slit position.}\\label{inicial} \\end{figure} ",
        "conclusions": "This work allows us to conclude that this system is a galaxy with polar ring, which has a companion galaxy with visible deformities, possibly due to the effects of the interaction process. A bridge or tails interconnect the PRG. The ring is corrugated and symmetrical. The host galaxy is elliptical and has some deformities."
    },
    "1208/1208.1193_arXiv.txt": {
        "abstract": "The compact radio source Sgr A* is coincident with a 4$\\times10^6$ \\msol black hole at the dynamical center of the Galaxy and is surrounded by dense orbiting ionized and molecular gas. We present high resolution radio continuum images of the central $3'$ and report a faint continuous linear structure centered on Sgr A* with a PA$\\sim60^0$. The extension of this feature appears to be terminated symmetrically by two linearly polarized structures at 8.4 GHz, $\\sim75''$ from Sgr A*. A number of weak blobs of radio emission with X-ray counterparts are detected along the axis of the linear structure. The linear structure is best characterized by a mildly relativistic jet from Sgr A* with an outflow rate 10$^{-6}$ \\msol\\, yr$^{-1}$.  The near and far-sides of the jet are interacting with orbiting ionized and molecular gas over the last 1--3 hundred years and are responsible for a 2$''$ hole, the ``minicavity\", characterized by disturbed kinematics, enhanced FeII/III line emission, and diffuse X-ray gas. The estimated kinetic luminosity of the outflow is $\\sim1.2\\times10^{41}$ erg s$^{-1}$, so the interaction with the bar may be responsible for the Galactic center X-ray flash inferred to be responsible for much of the fluorescent Fe K$\\alpha$ line emission from the inner 100pc of the Galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0195_arXiv.txt": {
        "abstract": "We present the chemical analysis of 49 giant stars of the globular cluster NGC~2419, using medium resolution spectra collected with the multi-object spectrograph DEIMOS@Keck. Previous analysis of this cluster revealed a large dispersion in the line strength of the infrared Ca~II triplet, suggesting an intrinsic star-to-star scatter in its Fe or Ca content. From our analysis, we assess that all the investigated stars share the same [Fe/H], [Ca/Fe] and [Ti/Fe] abundance ratios, while a large spread in Mg and K abundances is detected. The distribution of [Mg/Fe] is bimodal, with $\\sim$40\\% of the observed targets having subsolar [Mg/Fe], down to [Mg/Fe]$\\sim$--1 dex, a level of Mg-deficiency never observed before in globular clusters. It is found that the large dispersion in Mg abundances is likely the main origin of the observed dispersion of the Ca~II triplet lines strengths (that can be erroneously interpreted in terms of Fe or Ca abundance scatter) because Mg plays a relevant role in the atmosphere of giant stars as an electron donor. A strong depletion in the Mg abundance leads to an increase of the line strength of the Ca~II triplet, due to the variation in the electronic pressure, at a constant Fe and Ca abundance. Finally, we detect an anti-correlation between Mg and K abundances, not easily explainable within the framework of the current nucleosynthesis models. ",
        "introduction": "The old and metal-poor cluster NGC~2419 is by far the most luminous globular cluster (GC) residing in the outermost fringes of the Milky Way (MW) halo, $\\ga 10$ times brighter than any other cluster having $R_{GC}\\ge 40$~kpc, \\citep[see][for discussion and comparison with remote clusters in the halo of M31]{gal07}. This unusual feature, coupled with a half-light radius much larger than that of ordinary globulars of similar luminosity led several authors to the hypothesis that NGC~2419 can be the remnant of an originally larger system, like a nucleated dwarf galaxy \\citep{mvdb,ucd}.  Since the kinematics and the stellar mass function of NGC~2419 are fully consistent with a typical GC made up of stars and stellar remnants \\citep[see][and references therein]{iba11a,micml}, the key test to establish the actual nature of the system is to search for inhomogeneities in the abundance of chemical elements heavier than Aluminium among member stars. A spread in the heavy elements up to the iron-peak group would indicate that the system was a site of chemical evolution driven by Supernovae (SNe), implying that, at the epoch of SNe explosions, the progenitor of NGC~2419 was sufficiently massive to retain their highly energetic ejecta \\citep[$>$few$\\times 10^6~M_{\\sun}$;][]{baum}, hence it was likely a dwarf galaxy. On the other hand, GCs are observed to display virtually no spread in heavy elements and large (and correlated) spreads in light-elements (Na, O, Mg, Al, in particular), likely associated with a spread in He abundance \\citep[see][and references therein]{grat12}. These signatures are generally believed to trace early chemical evolution driven by much less energetic polluters than SNe, like massive Asymptotic Giant Branch stars \\citep[AGB][]{dant02,dercole08} or Fast Rotating Massive Stars \\citep[FRMS, see][and references therein]{dec07a,dec07b}, whose ejecta can be retained in systems in the plausible range of mass of proto-GCs. The anti-correlation between Na and O abundances is so ubiquitous in GCs (and non-existent in the field) that it has been proposed as a fundamental defining feature for globulars \\citep[with respect to open clusters and galaxies][]{carretta_def}.  The large distance of NGC~2419 \\citep[$\\sim$91 kpc][]{dicriscienzo11} has prevented the detailed spectroscopic analysis of a large sample of its stars, that is needed to perform this chemical test, until a couple of years ago\\footnote{\\citet{shetrone01} provided abundance analysis from high resolution spectroscopy for just {\\em one} star of NGC~2419.}. Recently \\citet[][C10 hereafter]{cohen10} presented the analysis of medium-resolution Keck-DEIMOS spectra around the Calcium triplet (near 8600 \\AA, CaT) for 43 Red Giant Branch (RGB) stars of NGC~2419. Using other GCs of known [Fe/H] as calibrators and adopting a constant [Ca/Fe] value they translated the measured pseudo- Equivalent Width of CaT and the magnitude difference between the stars and the cluster Horizontal Branch (HB, V-V$_{HB}$) into [Ca/H] values, as is usually done for [Fe/H] \\citep[see][and references therein]{rut,bat,starkenburg}. C10 found a significant spread in the [Ca/H] values derived with this method ([Ca/H]$_{CaT}$ hereafter): the distribution showed a strong peak at [Ca/H]$_{CaT}\\simeq -1.95$ and a long tail reaching [Ca/H]$_{CaT}\\simeq -1.45$. This relatively large spread in a heavy element \\citep[Fe or Ca, see][]{starkenburg} was interpreted by C10 as providing additional support to the hypothesis that NGC~2419 ``...is the remnant of a dwarf galaxy accreted long ago by the Milky Way...''. As a consistency check C10 obtained Fe and Ca abundances by spectral synthesis of Fe and Ca lines (other than CaT) that are present in their spectra. The puzzling result was that, while the Ca abundances are in reasonable agreement with those from CaT, no significant spread in Iron abundance was detected.  \\begin{figure} \\includegraphics[width=84mm]{hisfecat.ps} \\caption{Metallicity distribution derived from CaT, presented as ordinary histograms (lower panel) or as generalised histograms \\citep[a representation that removes the effects due to the choice of the starting point and of the bin width, see][]{laird}. Note that the distribution does not change if only the brightest stars are considered (dashed lines).} \\label{hisfecat} \\end{figure}  \\begin{figure} \\includegraphics[width=84mm]{cmd_CaT.ps} \\caption{Colour-Magnitude Diagram for the 118 stars included in the metallicity distribution from CaT shown in Fig.~\\ref{hisfecat} from the photometric catalog assembled by \\citet{b07} (see Section~\\ref{CaT}). Different symbols are adopted for different metallicity ranges.A grid of isochrones of different metallicity from the BaSTI set is overplotted \\citep{pietr04}, for reference.} \\label{cmd_CaT} \\end{figure}  Available photometry is not of much help in settling this issue, since in this low metallicity regime the sensitivity of optical colours to variations of metallicity is pretty weak. Moreover, even if a tiny colour spread is observed along the RGB, a spread in metallicity is not necessarily the only viable explanation. \\citet{dicriscienzo11b} used exquisite HST photometry to show that a small but significant colour spread is indeed detected at the base of the RGB, however this can be fully accounted for by the Helium spread that they invoke to reproduce the complex HB morphology of the cluster, within the classical scenario of multiple populations in GCs \\citep{grat12}.  Later \\citet[][C11 hereafter]{cohen11} used high-resolution Keck-HIRES spectra to derive the detailed abundance of several chemical elements in seven bright RGB members of NGC~2419. Interestingly, (a) the stars studied by C11 do not display any spread in [Fe/H] and [Ca/Fe] in excess to what is expected from the observational uncertainties, and (b) the abundance pattern of these stars is very similar to what is observed in the classical metal-poor GC  M30 (NGC~7099), residing in the inner halo of the MW ($D\\simeq 8$~kpc, $R_{GC}\\simeq 7$~kpc). One of the stars in the C11 sample (S1131), which is similar to the other six in all other aspects, was found to be extremely Mg-deficient ([Mg/Fe]=$-0.47$) and K-rich ([K/Fe]=$+1.13$), a strong anomaly never reported before, at least in Pop-II stars (see C11 for a detailed discussion). Hence, at present, the results of the chemical test on the nature of NGC~2419 are still not conclusive.  Here we report on the chemical analysis of the medium-resolution Keck-DEIMOS spectra obtained by \\citet[][I11a hereafter]{iba11a} and used, in that paper, for a thorough study of the kinematics of NGC~2419 \\citep[see also][]{iba11b}. ",
        "conclusions": "\\label{summa}  The main results derived from the analysis of 49 giant stars in NGC 2419 are summarised as follows:\\\\  \\begin{itemize}  \\item All the stars share the same iron content, with an average iron abundance of [Fe/H]=--2.09$\\pm$0.02 dex ($\\sigma$=~0.11 dex), where the observed dispersion is fully compatible with the uncertainties. Also, [Ca/Fe] and [Ti/Fe] turn out to be homogeneous;  \\item NGC 2419 exhibits a large dispersion in the Mg abundance, reaching values of [Mg/Fe]$\\sim$--1 dex (unusual for GC stars). The large spread of Mg is likely the main origin of the observed dispersion of the CaT lines strength, previously interpreted as intrinsic dispersion in Ca or Fe. In fact, a Mg depletion leads to an increase of the equivalent widths of the CaT lines (at a constant Ca abundance). This effect is confirmed by the fact that the iron content inferred from the Ca triplet lines clearly anti-correlates with the Mg abundances. However, we bear in mind that the strength of the CaT lines suffers from the abundances of all the elements that are electron donors and the correct line profile for these lines should be derived by taking into account the global budget of the free electrons;  \\item The [Mg/Fe] distribution is bimodal, with about 40\\% of the stars having sub-solar [Mg/Fe] abundance ratio. This fraction is similar to that suggested by \\citet{dicriscienzo11b} for the extreme population stars with initial helium abundance of Y$\\simeq$+0.4, according to the Horizontal Branch morphology of the cluster;  \\item a very large spread in K content is detected among the stars of NGC 2419, spanning from solar values up to [K/Fe]$\\sim$2 dex, with a bimodal distribution. A clearcut anti-correlation between [Mg/Fe] and [K/Fe] is observed, in agreement with the results already found by C11 for the star S1131, that shows an unusual depletion of Mg coupled with a strong enhancement of K.  \\end{itemize}  \\begin{figure} \\includegraphics[width=84mm]{naneky.ps} \\caption{[Mg/Fe] vs. [Fe/H] for {\\em a) open circles:} stars of NGC~2419 from the present study, {\\em b) triangles:} stars of NGC~2419 from C11, {\\em c) small dots:} stars in the dwarf spheroidal satellites of the MW from \\citet{kirby11}, {\\em d) $\\times$ symbols:} stars of various Galactic GCs from \\citet{carretta_uves}. Stars from NGC~7078 and NGC~2808 are plotted in blue and red, respectively, for reference, since they display the most Mg-deficient stars in the whole \\citet{carretta_uves} sample.} \\label{naneky} \\end{figure}  Fig.~\\ref{naneky} provides a direct illustration of the extremely unusual abundance pattern of NGC~2419. The spread in [Mg/Fe] is unrivalled both in GCs and in dwarf galaxies. A strong depletion of Mg is generally interpreted as due to a significant contribution of ejecta from SNIa to the gas from which Mg-poor stars are formed \\citep[see][and references therein]{bekki_lmc}. Indeed, Fig.~\\ref{naneky} shows that Mg-depletion with increasing metallicity is a common feature in dwarf spheroidal (dSph) galaxies and that [Mg/Fe]$<0.0$ stars are not rare in those systems. \\citet{carina} reports that the two most metal-rich stars in their sample of red giants in the Carina dSph have $-0.6<$[Mg/Fe]$<-1.0$, not far from the most Mg-deficient stars in our NGC~2419 sample. However, the Mg depletion in dSph (as in any other galaxies studied until now) is always coupled with an increase of [Fe/H]\\footnote{As well as a star formation history lasting for a few Gyrs \\citep{eline}, clearly not observed in NGC~2419 \\citep{dicriscienzo11b,micml}.}, since the whole effect is due to SNIa enriching the interstellar medium with material that is Fe-rich and poor in $\\alpha$ elements, thus reducing the [Mg/Fe] ratio with respect to the pattern previously set up by SNII \\citep[see][for a recent review]{eline}. This clearly does not occur in NGC~2419 where stars in the range $-1.4\\la$[Mg/Fe]$\\la+0.8$ are indistinguishable in terms of [Fe/H].  This fundamental observational fact, coupled with the analogy of the bi-modalities in [Mg/Fe] and HB morphology recalled above, suggest as more likely  the possibility that the observed pattern was produced by the same processes that cause the Na-O anti-correlations, and other signatures of the early chemical enrichment that are peculiar to GCs \\citep{carretta_def,grat12}. Further support to this hypothesis is provided by the fact that Mg-deficient stars lie systematically to the red of Mg-normal stars along the RGB, in the V, U-V Colour-Magnitude Diagram (Lardo et al., in preparation). In the framework of anti-correlations in GCs, Mg-deficient stars should correspond to Na-rich (and O-poor) stars and it is a generally observed characteristic of GCs to have Na-rich RGB stars that are redder than their Na-poor counterparts in CMDs that include the U-band \\citep[see][for discussion and references]{lardo}. The colour difference is due to variations in the strength of CN and CH features at wavelengths lower than $\\sim 4000$ \\AA, that are driven by variation of C and N abundance, that, in turn, are correlated with the abundance of other light elements tracing the self-enrichment process in GCs \\citep[like Na and O,][]{sbo11}.  Even if the observed Mg spread can be attributed to this kind of processes, the reason for the extreme behaviour of this cluster, as well as the origin of the Mg-K anti-correlation, remain unclear. This latter feature, if confirmed, may shed new light into the whole process of GC formation, possibly providing a tool that may help to discriminate between competing models of self-enrichment \\citep{dercole08,dec07a,dec07b}. An obvious singular characteristic of NGC~2419 is its distance from the centre of the Galaxy. Is it possible that a low degree of interaction with - for instance - the gas rich disc of the Milky Way has favoured the retention of enriched material in this cluster. \\citet{micml} found some indirect evidence for the cluster having a significant larger total mass in the past, as postulated by \\citet{dicriscienzo11b} on completely a different basis: also a large total stellar mass may have played a role (but this is a general feature of models of GC formation accounting for the presence of multiple populations). In any case, this mysterious cluster does not cease to reveal new interesting features that makes it especially worthy of further study."
    },
    "1208/1208.0640_arXiv.txt": {
        "abstract": "It is recently noted that solar eruptions can be associated with the contraction of coronal loops that are not involved in magnetic reconnection processes. In this paper, we investigate five coronal eruptions originating from four sigmoidal active regions, using high-cadence, high-resolution narrowband EUV images obtained by the Solar Dynamic Observatory (\\sat{sdo}). The magnitudes of the flares associated with the eruptions range from the \\sat{goes}-class B to X. Owing to the high-sensitivity and broad temperature coverage of the Atmospheric Imaging Assembly (AIA) onboard \\sat{sdo}, we are able to identify both the contracting and erupting components of the eruptions: the former is observed in cold AIA channels as the contracting coronal loops overlying the elbows of the sigmoid, and the latter is preferentially observed in warm/hot AIA channels as an expanding bubble originating from the center of the sigmoid. The initiation of eruption always precedes the contraction, and in the energetically mild events (B and C flares), it also precedes the increase in \\sat{goes} soft X-ray fluxes. In the more energetic events, the eruption is simultaneous with the impulsive phase of the nonthermal hard X-ray emission. These observations confirm the loop contraction as an integrated process in eruptions with partially opened arcades. The consequence of contraction is a new equilibrium with reduced magnetic energy, as the contracting loops never regain their original positions. The contracting process is a direct consequence of flare energy release, as evidenced by the strong correlation of the maximal contracting speed, and strong anti-correlation of the time delay of contraction relative to expansion, with the peak soft X-ray flux. This is also implied by the relationship between contraction and expansion, i.e., their timing and speed. ",
        "introduction": "It is generally regarded that solar eruptions are due to a disruption of the force balance between the upward magnetic pressure force and the downward magnetic tension force. Since the eruption can only derive its energy from the free energy stored in the coronal magnetic field \\citep{forbes00}, ``the coronal field lines must contract in such a way as to reduce the magnetic energy $\\int_\\mathcal{V} B^2/8\\pi$'' \\citep{hudson00}. The contraction must be associated with the reduction of the magnetic tension force for each individual loop-like field line undergoing contraction, as its footpoints are effectively anchored in the photosphere. Eventually a new force balance would be achieved between the magnetic pressure and tension force after the energy release. From an alternative viewpoint, the average magnetic pressure $B^2/8\\pi$ must decrease over the relevant volume $\\mathcal{V}$ across the time duration of the eruption. $\\mathcal{V}$ can be roughly regarded as the flaring region, primarily in which magnetic energy is converted into other forms of energies. The contraction process, termed as ``magnetic implosion'' by \\citet{hudson00}, is very similar to the shrinkage of post-flare loops \\citep{fa96}, except that loop shrinkage is driven by temporarily enhanced magnetic tension force at the cusp of the newly reconnected field lines, whereas loop contraction by reduced magnetic pressure in the flaring region. Additionally, with newly reconnected loops piling up above older ones, the post-flare arcade as a whole often expands rather than shrinks with time.  \\citet{hudson00} concluded that ``a magnetic implosion must occur simultaneously with the energy release'' , based on no assumption about the energy release process itself. However, the detailed timing and location of loop contraction might provide diagnostic information on the eruption mechanism. For example, when the reconnection-favorable flux emerges inside a filament channel \\citep[Figure \\ref{model}(a); adapted from][]{cs00}, it cancels the small magnetic loops below the flux rope, which results in a decrease of the local magnetic pressure. The whole dipolar magnetic structure must contract correspondingly. Meanwhile, plasmas on both sides of the polarity-inversion line (PIL) would move inward to form a current sheet below the flux rope and the subsequent evolution could follow the paradigm of the standard flare model \\citep[e.g.,][]{kp76}. In that case, overlying coronal loops could be observed to initially contract and then erupt. In a different scenario, a twisted flux rope confined by potential-like magnetic fields is found to be energetically favorable to ``rupture'' through the overlying arcade via ideal-MHD processes \\citep[Figure \\ref{model}(b); adapted from ][]{sturrock01}. This is clearly demonstrated in MHD simulations by \\citet{gf06} and \\citet{rachmeler09}, in which overlying loops can be seen to be pushed upward and aside as the flux rope kinks and expands, and after the rope ruptures through the arcade, overlying loops on both sides quickly contract toward the core region, due to the reduction of the magnetic pressure in the core field with the escape of the flux rope. In particular for this scenario (\\fig{model}(b)), one would expect to see both the expanding flux rope and the contracting overlying loops during the eruption as long as the arcade is only partially opened. Although both scenarios involve a pre-existent flux rope, supposedly they can also accommodate those models in which the flux rope forms immediately prior to \\citep[e.g.,][]{moore01}, or during the course of \\citep[e.g.,][]{adk99}, the eruption.    Corresponding to the aforementioned models (\\fig{model}), our previous observational studies also suggest two different scenarios, i.e., 1) the bunch of coronal loops undergoing contraction later becomes the front of the eruptive structure \\citep{lwa09}; and 2) the contracting loops are distinct from the eruptive structure \\citep{lw09, lw10}. The role of contraction in the eruption, however, has been unclear in both scenarios. For Scenario 1, the event reported by \\citet{lwa09} remains unique in the literature; as for Scenario 2, the eruptive structure is not easy to detect before its appearance as a CME in coronagraph, unless there is dense filament material serving as the tracer \\citep{lw09}. In some cases its slow ascension and expansion during the early stage might manifest as the gradual inflation of overlying coronal loops \\citep{liu10a}. Only with the advent of the Solar Dynamic Observatory \\citep[\\textit{SDO};][]{pesnell12} which provides a continuous and wide temperature coverage, is the eruptive structure itself more frequently identified beneath the coronagraph height as a hot, diffuse plasmoid \\citep[e.g.,][]{liu10b, cheng11}.  Here in a further investigation of Scenario 2, we identify both the erupting and contracting components using \\sat{sdo} data, hence for the first time we are able to study in detail their relationship as well as the implication for the eruption mechanism and the associated energy release process.  In the rest of the paper, we present in Section 2 the results of the investigation on five flares (Table 1) observed by the Atmospheric Imaging Assembly \\citep[AIA;][]{lemen12} onboard \\sat{sdo}, and we make concluding statements in Section 3.     \\begin{figure}\\epsscale{0.8}% \\plotone{model.eps} \\caption{CME models relevant to magnetic implosion. \\textit{a}) Schematic diagram of the emerging flux triggering mechanism for CMEs \\citep[adapted from][]{cs00}. The emerging flux inside the filament channel cancels the pre-existing loops, which results in the in-situ decrease of the magnetic pressure. Magnetized plasmas are driven inward to form a current sheet beneath the flux rope; \\textit{b}) Schematic sketch showing that in the three-dimensional space a twisted flux rope can rupture the overlying magnetic arcade and erupt by pushing the magnetic arcade aside \\citep[adapted from][]{sturrock01}. With the escape of the flux rope, the arcade field undergoes a contraction due to the decreased magnetic pressure in the core field. \\label{model}} \\end{figure}  \\begin{figure}\\epsscale{0.85} \\plotone{sigmoid.eps} \\caption{Pre-flare configuration for the five flares studied. Left column: line-of-sight magnetograms obtained by the Helioseismic and Magnetic Imager (HMI) onboard \\sat{sdo}; Middle and Right columns: corresponding EUV images in the cold and warm/hot AIA channels, respectively, showing the sigmoidal morphology and structure. For AR 11158 (3rd and 4th rows), we use HMI vector magnetograms to construct nonlinear force-free field (NLFFF; see the text for details). The extrapolated field lines are color-coded according to the intensity of vertical currents on the surface. \\label{sigmoid}} \\end{figure} ",
        "conclusions": "\\begin{figure}\\epsscale{0.8} \\plotone{correl.eps} \\caption{Correlation of the maximal contraction/expansion speed, $V$, and the time delay of contraction relative to expansion, $\\Delta t$, with the flare magnitude in terms of the peak \\sat{goes} 1--8 \\AA\\ flux. The confidence level of the linear correlation coefficient, $cc$, of $\\lg(V)$ and $\\lg \\Delta t$ with $\\lg(F)$ is given in the brackets. \\label{correl}} \\end{figure}  We have investigated four sigmoidal active regions, in which five eruptions with signatures of magnetic implosion occurred. The magnitudes of the flares associated with the eruptions span almost the whole flare ``spectrum'', from the \\sat{goes}-class B to X. In all of the flares studied, there are both a contracting and an erupting component: the former is only observed in cold EUV channels and the latter is preferentially visible in warm/hot EUV channels. This is because the contracting component is composed of large-scale, potential-like coronal loops overlying the elbows of the sigmoid, while the erupting component is associated with newly reconnected flux tubes originating from the center of the sigmoid \\citep[c.f.,][]{liu10b, aulanier10, schrijver11}. Several important aspects of these observations are discussed as follows.  \\begin{itemize} \\item \\textbf{Consequence of loop contraction:} the overlying loops undergoing contraction never regain their pre-flare positions, which implies a new equilibrium with reduced magnetic energy as the eruption is powered by magnetic energy. One may argue that the apparent contraction of coronal loops could be a projection effect, i.e., the loop plane tilts due to the flare impulse. But in that case, one would expect the restoration of the loops once the flare impulse has passed away. In observation, however, the contracting loops may oscillate about a lower height \\citep[e.g., \\fig{ltc0213}; see also][]{lw10}, but never reach the original heights after the eruption. Thus, the contraction within the loop plane must make a significant contribution.  \\item \\textbf{Correlation between contraction and eruption:} the contraction speed seems to depend on the intensity/magnitude of the eruption. From \\fig{correl}, one can see that despite this very small sample size, the peak \\sat{goes} SXR flux as a proxy of the flare magnitude is linearly correlated very well with the measured maximal contracting speed in the log-log plot, although not so well with the maximal erupting speed. Unlike contracting loops which are clearly defined, however, the measurement of the erupting speed involves larger uncertainties as the front of the expanding bubble tend to get more and more diluted and eventually overwhelmed by the background during propagation, thereby leading to underestimation of its speed. One more caveat to keep in mind is that these speeds are not necessarily measured at the time of the peak SXR flux.   \\item \\textbf{Timing:} the eruption precedes the contraction in all of the flares studied, which establishes loop contraction as a consequence of eruption. There is also a trend that the more energetic the eruption, the smaller the time delay of the loop contraction relative to the onset of the expansion of the erupting component, which is demonstrated in \\fig{correl} as a strong anti-correlation between the time delay and the peak \\sat{goes} SXR flux in the log-log plot. This time delay is presumably determined by the expansion speed of the erupting component. In addition, in the relatively weak B- and C-flares, the initiation of the erupting component precedes the increase in \\sat{goes} SXR fluxes; but in the stronger M- and X-flares, it is concurrent with the increase in nonthermal HXR fluxes. This may lend support to \\citet{lin04}, who concluded that CMEs are better correlated with flares if there is more free energy available to drive the eruption. On the other hand, since the CME progenitor, i.e., the expanding bubble, forms before the flare onset as the weak events clearly demonstrate, the CME must be independent of the conventionally defined flare, or, the flare is only a byproduct of the CME, unless the eruption mechanism for the weak events is different from that for the energetic ones.  \\item \\textbf{Asymmetry of contraction:} the two groups of coronal loops overlying the elbows of the sigmoid often contract asymmetrically, i.e., not only they contract at different speeds but either group could show little sign of contraction, dependent on the detailed interaction between the core field and the arcade field, including, presumably, their relative strength and the spatial distribution of the decay index of the restraining field \\citep{kt06, lag09, liuc10}. For the 2010 August 1 event in particular, \\citet{liu10b} concluded that the majority of the flare loops were formed by reconnection of the stretched legs of the less sheared loops overlying the southern elbow and the center of the sigmoid, based on the reconnection rate inferred from the H$\\alpha$ ribbon motion. The eruption therefore left most loops overlying the northern elbow unopened. This explains why only these loops underwent obvious contraction. The intensity/magnitude of the eruption could be another relevant factor as among the events studied only those greater than M-class show contraction of loops overlying both elbows of the sigmoid.  \\item \\textbf{Implication for eruption mechanism:} as the contracting component is distinct from the erupting component, we conclude that these eruptions conform to the ``rupture model'' in which the arcade field is partially opened \\citep[][\\fig{model}(b)]{sturrock01}. We can further exclude the breakout model because the coronal loops undergoing contraction are arched over, rather than located to the side of, the sheared core field. The loop contraction in the latter occasion results from reconnection at the magnetic null above the central lobe in the breakout model. \\end{itemize}  In conclusion, these observations substantiate the loop contraction as an integrated process in eruptions of sigmoidal active regions in which the restraining arcade field is only partially opened, consistent with theoretical expectations. The consequence of loop contraction is a new equilibrium of the coronal field with reduced magnetic energy, and the process itself is a result of the flare energy release, as evidenced by the strong correlation of the maximal contracting speed, and strong anti-correlation of the time delay of contraction relative to expansion, with the peak SXR flux."
    },
    "1208/1208.4512_arXiv.txt": {
        "abstract": "{Isolated starless cores within molecular clouds can be used as a testbed to investigate the conditions prior to the onset of fragmentation and gravitational proto-stellar collapse.} {We aim to determine the distribution of the dust temperature and the density of the starless core B68.} {In the framework of the \\textit{Herschel} Guaranteed-Time Key Programme ``The Earliest Phases of Star formation'' (EPoS), we have imaged B68 between 100 and $500~\\mu$m. Ancillary data at (sub)millimetre wavelengths, spectral line maps of the \\element[][12]{CO}~(2--1), and \\element[][13]{CO}~(2--1) transitions as well as an NIR extinction map were added to the analysis. We employed a ray-tracing algorithm to derive the 2D mid-plane dust temperature and volume density distribution without suffering from the line-of-sight averaging effects of simple SED fitting procedures. Additional 3D radiative transfer calculations were employed to investigate the connection between the external irradiation and the peculiar crescent-shaped morphology found in the FIR maps.} {For the first time, we spatially resolve the dust temperature and density distribution of B68, convolved to a beam size of 36\\farcs4. We find a temperature gradient dropping from $(16.7 {+1.3 \\above0pt -1.0})$~K at the edge to $(8.2 {+2.1 \\above0pt -0.7})$~K in the centre, which is about $4$~K lower than the result of the simple SED fitting approach. The column density peaks at $N_\\mathrm{H} = (4.3 {+1.4 \\above0pt -2.8})~\\times 10^{22}$~cm$^{-2}$, and the central volume density was determined to $n_\\mathrm{H} = (3.4 {+0.9 \\above0pt -2.5})~\\times 10^{5}$~cm$^{-3}$. B68 has a mass of $3.1~M_{\\sun}$ of material with $A_K > 0.2$~mag for an assumed distance of 150~pc. We detect a compact source in the southeastern trunk, which is also seen in extinction and CO. At 100 and $160~\\mu$m, we observe a crescent of enhanced emission to the south.} {The dust temperature profile of B68 agrees well with previous estimates. We find the radial density distribution from the edge of the inner plateau outward to be $n_\\mathrm{H} \\propto r^{-3.5}$. Such a steep profile can arise from either or both of the following: external irradiation with a significant UV contribution or the fragmentation of filamentary structures. Our 3D radiative transfer model of an externally irradiated core by an anisotropic ISRF reproduces the crescent morphology seen at 100 and 160~$\\mu$m. Our CO observations show that B68 is part of a chain of globules in both space and velocity, which may indicate that it was once part of a filament that dispersed. We also resolve a new compact source in the southeastern trunk and find that it is slightly shifted in centroid velocity from B68, lending qualitative support to core collision scenarios.} ",
        "introduction": "There is a general consensus about the formation of low and intermediate-mass stars that they form via gravitational collapse of cold pre-stellar cores. A reasonably robust evolutionary sequence has been established over the past years from gravitationally bound pre-stellar cores \\citep{ward-thompson02} over collapsing Class~0 and Class~I protostars to Class~II and Class~III pre-main sequence stars \\citep{andre93,shu87}. The scenario of inside-out collapse of a singular isothermal sphere \\citep[e.g.][]{shu77,young05} seems to be consistent with many observations of cores that already host a first hydrostatically stable protostellar object \\citep[e.g.][]{motte01}. There are also disagreeing results in the case of starless cores, where inward motions exist before the formation of a central luminosity source \\citep{difrancesco07}, and infall can be observed throughout the cloud, not only in the centre \\citep[e.g.][]{tafalla98}.  \\begin{table*}[t] \\caption{Details of the \\textit{Herschel} observations of B68.} \\label{t:obs} \\begin{tabular}{ccccccccc} \\hline\\hline OD       & OBSID & Instrument & Wavelength & Map size & Repetitions & Duration & Scan speed & Orientation \\\\ & & & ($\\mu$m) & & & (s) & (\\arcsec~s$^{-1}$) & (\\degr) \\\\ \\hline 287      & 1342191191 & SPIRE & 250, 350, 500 & $9\\arcmin\\times9\\arcmin$ & 1 & 555 & 30 & $\\pm$ 42\\\\ \\raisebox{-5pt}[0pt][-5pt]{320} & 1342193055 & \\raisebox{-5pt}[0pt][-5pt]{PACS} & \\raisebox{-5pt}[0pt][-5pt]{100, 160} & \\raisebox{-5pt}[0pt][-5pt]{$7\\arcmin\\times7\\arcmin$} & 30 & 4689 & 20 & 45 \\\\ & 1342193056 & & & & 30 & 4568 & 20 & 135 \\\\ \\hline \\end{tabular} \\tablefoot{The orientation angles of the scan direction listed in the last column are defined relative to the instrument reference frame. The PACS maps were obtained using the \\textit{homogeneous coverage} and \\textit{square map} options of the scan map AOT. The SPIRE maps were obtained separately.} \\end{table*}  \\citet{evans01} and \\citet{andre04} consider positive temperature gradients, which are caused by external heating and internal shielding, for their density profile models of pre-stellar cores. Although their results differed quantitatively, both come to the conclusion that the slopes of the density profiles in the centre of pre-stellar cores are significantly smaller than previously assumed.  However, these temperature profiles were derived from radiative transfer models with certain assumptions about the local interstellar radiation field, but lack an observational confirmation.  To improve this situation and measure the dust temperature and density distribution in a realistic way, we used the imaging capabilities of the Photodetector Array Camera and Spectrograph (PACS\\footnote{PACS has been developed by a consortium of institutes led by MPE (Germany) and including UVIE (Austria); KU Leuven, CSL, IMEC (Belgium); CEA, LAM (France); MPIA (Germany); INAF-IFSI/OAA/OAP/OAT, LENS, SISSA (Italy); IAC (Spain). This development has been supported by the funding agencies BMVIT (Austria), ESA-PRODEX (Belgium), CEA/CNES (France), DLR (Germany), ASI/INAF (Italy), and CICYT/MCYT (Spain).}) \\citep{pacs} and the Spectral and Photometric Imaging Receiver (SPIRE) \\citep{spire} instruments on board the \\textit{Herschel} Space Observatory \\citep{herschel} to observe the Bok globule \\object{Barnard~68} (B68, LDN~57, CB~82) \\citep{barnard27} as part of the \\textit{Herschel} Guaranteed Time Key Programme ``The Earliest Phases of Star formation'' \\citep[EPoS; P.I.~O.~Krause; e.g.][]{beuther10,beuther12,henning10,linz10,stutz10,ragan12} of the PACS consortium. With these observations, we close the important gap that covers the peak of the spectral energy distribution (SED) that determines the temperature of the dust. These observations are complemented by data that cover the SED at longer wavelengths reaching into the millimetre range.  Bok globules \\citep{bok47} like B68 are nearby, isolated, and largely spherical molecular clouds that in general show a relatively simple structure with one or two cores \\citep[e.g.~][]{larson72,keene83,chen07,stutz10}. A number of globules appear to be gravitationally stable and show no star formation activity at all. Whether or not this is only a transitional stage depends on the conditions inside and around the individual object \\citep[e.g.~][]{launhardt10} like its temperature and density structure, as well as external pressure and heating. Irrespective of the final fate of the cores and the globules that host them, such sources are ideal laboratories in which to study physical properties and processes that exist prior to the onset of star formation.  B68 is regarded as an example of a prototypical starless core. It is located at a distance of about 150~pc \\citep{degeus89,hotzel02co,lombardi06,alves07} within the constellation of Ophiuchus and on the outskirts of the \\object{Pipe nebula}. The mass of B68 was estimated to lie between 0.7 \\citep{hotzel02co} and $2.1~M_{\\sun}$ \\citep{alves01nature}, which also relies on different assumptions for the distance. Previous temperature estimates were mainly based on molecular line observations that trace the gas and resulted in mean values of the entire cloud between 8~K and 16~K \\citep{bourke95,hotzel02temp,bianchi03,lada03,bergin06}. Dust continuum temperature estimates lie in the range of 10~K \\citep{kirk07}. However, they usually lack the necessary spatial resolution and spectral coverage for a precise assessment of the dust temperature. Typically, the Rayleigh-Jeans tail is well constrained by submillimetre observations. However, the peak of the SED, which is crucial for determining the temperature, is measured either with little precision or too coarse a spatial resolution that is unable to sample the interior temperature structure.  The radial density profile of B68 is often represented by fitting a Bonnor-Ebert sphere (BES) \\citep{ebert55,bonnor56}, where self gravity and external and internal pressure are in equilibrium. Based on deep near-infrared (NIR) extinction mapping, \\citet{alves01nature} show by using this scheme that B68 appears to be on the verge of collapse. However, this interpretation is based on several simplifying assumptions, e.g.~spherical symmetry and isothermality that is supposed to be caused by an isotropic external radiation field. \\citet{pavlyuchenkov07} have shown that even small deviations from the isothermal assumption can significantly affect the interpretation of the chemistry derived from spectral line observations, especially for chemically almost pristine prestellar cores like B68, and therefore also affect the balance of cooling and heating.  We demonstrate in this paper that none of these three conditions are met, and we conclude that even though the radial column density distribution can be fitted by a Bonnor-Ebert (BE) profile, any interpretation of the nature of B68 that is based on this assumption should be treated with caution. ",
        "conclusions": "\\label{s:conclusions} We have presented FIR continuum maps of the isolated low-mass prestellar core, B68, obtained with \\textit{Herschel}. Using ancillary sub-mm continuum maps, all convolved to the resolution of our SPIRE $500~\\mu$m data, we presented two approaches to modelling the dust temperature and column density distribution: line-of-sight-averaged SED fitting and modelling by ray-tracing. The latter modelling approach permitted us to also calculate the volume density distribution.  We found the dust temperature rises from 8.2~K in the centre of 16.0~K at the edge. The tenuous medium to the southeast, facing the galactic plane, reaches almost 20~K. The column density peaks at $N_\\mathrm{H} = 4.3 \\times 10^{22}~\\mathrm{cm}^{-2}$ and declines to $8 \\times 10^{20}~\\mathrm{cm}^{-2}$. B68 has a total mass of $4.2~M_{\\sun}$, or $3.1~M_{\\sun}$ when considering only material with $A_K > 0.2$~mag. We detected a new compact source in the southeastern trunk of B68 at 160 and $250~\\mu$m, which has a mass of about $0.08~M_{\\sun}$.  The mean radial column density profile can be fitted by a power-law with $N_\\mathrm{H} \\propto r^{-2.5}$ between the flat inner core to the edge. The corresponding volume density profile follows $n_\\mathrm{H} \\propto r^{-3.5}$, ranging from $3.4 \\times 10^5~\\mathrm{cm}^{-3}$ (centre) to $4 \\times 10^2~\\mathrm{cm}^{-3}$ (edge), which is consistent with a \"Plummer-like\" profile with $\\eta = 4$ \\citep{whitworth01}. However, we show that the mean radial profile is not a good representation of the density distribution beyond $40''$ (6000~AU), because deviations of up to a factor of 4 are seen in the volume density profile beyond this radius. We will address this in detail with 3D modelling in a forthcoming work.  We accounted for the effects of data-handling and differing angular resolutions in our analysis of the core profiles. We also explored two physical scenarios that can, either individually or in combination, explain the density profile: (i) an isolated core in thermal and mechanical equilibrium irradiated by an external ISRF with a significant UV component \\citep[e.g.][]{andre00,whitworth01}, and (ii) the remnants of a non-magnetic isothermal cylinder \\citep{ostriker64}. Both scenarios can produce density profiles that are in better agreement with our results, which is steeper than the commonly assumed singular isothermal sphere \\citep[$n_\\mathrm{H} \\propto r^{-2}$,][]{shu87}.  We observed a unique crescent-shaped asymmetry in the 100 and $160~\\mu$m emission profiles, intriguingly similar to the morphology shown in C$^{18}$O by \\citet{bergin02}. Using 3D radiative transfer modelling, such a morphology can be reproduced in the continuum by an anisotropically irradiated core, including contributions from the nearby B2IV star $\\theta$~Oph and from the Galactic centre and galactic plane ISRF. Put together, this demonstrates the importance of accounting for an anisotropic ISRF when modelling a globule in detail.  \\begin{figure} \\centering \\resizebox{0.789\\hsize}{!}{\\includegraphics{B68_13CO_radvel}} \\caption{\\label{f:13corvel}Radial velocity field of the larger area around B68 seen in the \\element[][13]{CO}~(2--1) line. There appears to be a gradient with increasing radial velocities from north to south, which is consistent with a gas flow of constant velocity that is inclined towards the LoS.} \\end{figure}  We presented observations of $^{12}$CO and $^{13}$CO of a larger area surrounding B68 than in previous molecular line studies. The new compact source in the southeastern trunk is the peak in the integrated intensity, and it also exhibits enhanced linewidths ($\\sim 1.1$ and $0.65~\\mathrm{km~s}^{-1}$, for $^{12}$CO and $^{13}$CO, respectively) and a slightly redshifted centroid velocity. Our observations lend qualitative credence to the ``bullet'' scenario posited by \\citet{burkert09}, in which B68 collides with a small core (our new, kinematically distinct compact source), triggering gravitational collapse.  In addition, neighbouring globules, B69 and B71, show a smooth velocity gradient, suggesting that this system of globules were once part of a single filament that dispersed. This agrees with the predictions of the fragmenting isothermal cylinder scenario mentioned above."
    },
    "1208/1208.4038_arXiv.txt": {
        "abstract": "The structure of the deflagration burning front in type Ia supernovae is  considered. The parameters of the flame are obtained: its normal velocity and thickness. The results are in good agreement with the previous works of different authors. The problem of pulsation instability of the flame, subject to plane perturbations, is studied. First, with the artificial system with switched--off hydrodynamics the possibility of secondary reactions to stabilize the front is shown. Second, with account of hydrodynamics, realistic EOS and thermal conduction we can obtain pulsations when Zeldovich number was artificially increased. The critical Zeldovich numbers are presented. These results show the stability of the flame in type Ia supernovae against pulsations because its effective Zeldovich number is small. ",
        "introduction": "Supernovae explosions are among  the most spectacular  events in the Universe: their energy release significantly mixes the interstellar medium and acts like a driving force in gas dynamics of galaxies and production of cosmic rays. The luminosity of an exploding star becomes comparable with the luminosity of the progenitor galaxy, and allows to observe processes in the most distant regions of the Universe. If regular features of the supernova explosions are found for any subtype of supernovae, it could open a new way to measure cosmological distances and the values of the cosmological parameters.  Despite a long history of investigations of these events the complete understanding of underlying physics is still missing. There are several kind of supernova explosions of different types of the progenitor stars with absolutely different physical phenomena behind them. Here we will consider only one subtype of explosions, namely the thermonuclear explosions. These kind of supernovae is called the  supernovae of type Ia, SNIa. Analysis of the observational data indicates that the explosion is induced  by the thermonuclear burning of premixed carbon--oxygen fuel. Such phenomenon usually takes place in degenerate stars, white dwarfs.   The mode of the explosive nuclear burning in supernovae is still a controversial issue, in spite of many years of the research in the field. Four decades ago, \\citet{Arn} was the first to consider supersonic combustion, i.e. detonation, in supernovae. Later, \\citet{IIC} obtained a sub-sonic flame (deflagration) propagating in spontaneous regime with pulsations and a subsequent transition to detonation, while \\citet{NSN} considered the deflagration propagating due to convective heat transfer. Both detonation and deflagration have their merits and problems in explaining the supernova phenomenon (see, e.g. \\cite{WW}). It is not clear if detonation succeeds or fails to develop, but it is clear that in any case the combustion must be much faster than it is suggested by the analysis of the propagation of a laminar one-dimensional flame.  From microscopic point of view one-dimensional nuclear flame is a wave described  essentially in the same way as it was done by \\citet{ZFK} in spite of complications introduced by nuclear kinetics and very high conductivity of dense presupernova matter. It is found that the conductive flame propagates in  presupernova with the speed which is too slow to explain the supernova outburst correctly since the flame Mach number is of the order of one percent or less \\citep{TimmesWoosley_ApJ_1992}.   The fuel consumption can naturally be accelerated by the development of the instabilities inherent to the flame front. As it is explained in the classical paper by \\citet{Landau,Landau2}, the hydrodynamic instability leads to wrinkling or roughening of the front surface, and hence to an increase of its area with respect to the smooth front and consequently to an acceleration of the flame propagation. In some cases, observed in laboratory experiments like \\cite{Gostintsev}, when the LD instability is really strong (large density jump accross the flame front, and the flame is freely expanding) such instabilities can lead to a transition from the regime of slow flame propagation  to the  regime of detonation. Since the flame propagates in gravitational field, and the burned ashes have lower density than the unburned fuel, the Rayleigh--Taylor (RT) instability is often considered to be the dominant instability governing the corrugation of the front \\cite{MA,wooa,woob,LivA,kho}. The RT instability creates turbulent cascade providing  an acceleration of the flame front. However it leads to additional difficulties in modeling the SN event \\cite{NHt,Niem,Nwoo,WoosleyKersteinEtAl_ApJ_2010,AspdenEtAl_ApJ_2008}.   It is well known \\citep{Landau,Landau2,LL,Will} % that large portions of a slow planar flame front are unstable with respect to the large scale bending. This universal instability is called the Landau-Darrieus (LD) instability. For the wavelengths much longer than the flame thickness it does not depend, on complex processes which take place in the burning zone. Development of the LD instability depends only on the sign of $\\Delta\\rho = \\rho_{\\rm u} -\\rho_{\\rm b}$, where $\\rho_{\\rm u}$ and $\\rho_{\\rm b}$ are the densities of the unburned and burned ``gases'' respectively. The LD instability of the planar flame fronts with respect to large scale bending takes place if and only if $\\Delta\\rho >0$.  The LD instability plays an important role in many physical phenomena such as the usual chemical burning of gases, explosive boiling of liquids \\citep{Frost}, electroweak phase transitions \\citep{KamFr}, dynamics of thermally bistable gas \\citep{AMS2}, and thermonuclear burning in supernovae  \\citep{BSW,NHld}. The detailed consideration the of non-linear stage of the LD instability and the calculation of the fractal dimension of the flame front for this case is given by % \\citet{BS,Joulin_PRE_1994}. It is interesting, and somewhat puzzling, that a similar dependence of the flame fractal dimension on the density discontinuity was found in the 3D SPH simulations of the flame subject to the RT instability \\citep{BrGar}.  Both LD and RT instabilities develop on scales much larger than the flame thickness and they can be successfully studied in the approximation of the discontinuous front. This approximation is not valid for another instability, first discovered by \\citet{Zelpw} % in his investigation of the powder combustion. This instability originated from a strong temperature dependence of the reactions rates. As a result of  that local fluctuations of the heating rate caused by the temperature fluctuations cannot be controlled by  thermal conduction. This phenomenon can lead to a pulsating regime of the front propagation and to a renormalization of the mean front velocity \\cite{W}. Such instability can develop even for one-dimensional perturbations when the plane front preserves its shape. We denote it as TP (thermal-pulsational) instability. After publication of paper by \\citet{BarZI} TP instability was studied quantitatively in many works (for the list of references see the book \\cite{Zel}).  A very nice review on the SNIa physics is given in \\cite{HillebrandtNiemeyer_astro_ph_0006305} (see also \\cite{RopkeHillebrandtBlinnikov2006,RopkeHillebrandtEtAl_ApJ_2007}). There exist several scenarios of SNIa explosions. The most popular are the following: the single--degenerate scenario, the double--degenerate scenario, and the sub-Chandrasekhar mass explosion. In this paper the single degenerate scenario is considered.  The paper is organized as follows. In Section \\ref{sec:model} a model of a white dwarf explosion is presented and the physical conditions are discussed. In  Section \\ref{sec:flame_prop_1d} the stationary propagating flame is considered analytically and numerically. The dependence of the results on nuclear reaction network is discussed. In Section \\ref{sec:th_dyn_ins} artificial systems with switched off hydrodynamics are considered. The effect of secondary reaction on pusations is considered. In Section \\ref{sec:front_puls} the stability of the flame under conditions close to those in a white dwarf is considered. We show that pulsations could exist in this system when the Zeldovich number is aftificially increased. After it we make conclusions about stability of real flames. ",
        "conclusions": "One-dimensional flame propagation in presupernova white dwarf has been considered. Flame properties for different star densities were obtained. It is shown that when only one nuclear reaction in the nuclear network is considered, the flame velocity strongly differs from more sophisticated net simulations. So this simplified approach can be used only in approximate simulations when the exact value of the flame velocity is not required, or alternatively the nuclear rate should be adjusted to fit the correct value of $v_n$.  We have also studied one-dimensional pulsational instability. First, with artificial systems the switched--off hydrodynamics were considered. The possibility of secondary reactions to stabilize the front was shown with the help of this system (by decreasing the net {\\sf Ze} number).  In ``real'' simulations, presented in Section \\ref{sec:flame_prop_1d} no pulsation regime was found, so we used the simplified Arrhenius law in one-step reaction to make it possible to change Zeldovich number, {\\sf Ze}. For high {\\sf Ze} numbers the pulsations do exist. It means that our numerical code can resolve such pulsations and that they can exist in conditions close to ``real'', but with steeper energy generation rate. By means of numerical simulations we have obtained critical Zeldovich numbers for densities in the range $\\rho=[2\\cdot 10^8, 2\\cdot 10^9]$ g/cm$^3$. These values are larger than those in real flame, which proves one-dimensional stability of the realistic flame fronts in supernovae.  \\bigskip \\centerline{\\bf Acknowledgements} We are grateful to V.Chechetkin, W.Hillebrandt, A.Kruzhilin, J.Niemeyer, P.Sasorov, S.Woosley, and F.Timmes for cooperation, discussions and references. MPA supported the work of SB in Germany. The work in Russia is supported by RFBR grants 11-02-00441-a and 13-02-92119, Sci. Schools 5440.2012.2, 3205.2012.2 and 3172.2012.2, by the contract No.~11.G34.31.0047 of the Ministry of Education and Science of the Russian Federation, and SCOPES project No.~IZ73Z0-128180/1."
    },
    "1208/1208.0331_arXiv.txt": {
        "abstract": "In this letter, 21\\,cm intensity maps acquired at the Green Bank Telescope are cross-correlated with large-scale structure traced by galaxies in the WiggleZ Dark Energy Survey.  The data span the redshift range $0.6 < z < 1$ over two fields totaling $\\sim41$~deg.~sq. and $190$~hr of radio integration time.  The cross-correlation constrains $\\Omega_{\\rHI} b_{\\rHI}r=[0.43\\pm0.07({\\rm stat.})\\pm0.04({\\rm sys.})]\\times 10^{-3}$, where $\\Omega_{\\rHI}$ is the neutral hydrogen (\\HI) fraction, $r$ is the galaxy--hydrogen correlation coefficient, and $b_{\\rHI}$ is the \\HI\\ bias parameter.  This is the most precise constraint on neutral hydrogen density fluctuations in a challenging redshift range.  Our measurement improves the previous 21\\,cm cross-correlation at $z\\sim0.8$ both in its precision and in the range of scales probed. ",
        "introduction": "Measurements of neutral hydrogen are essential to our understanding of the universe.  Following cosmological reionization at $z\\sim6$, the majority of hydrogen outside of galaxies is ionized.  Within galaxies, it must pass through its neutral phase (\\HI) as it cools and collapses to form stars. The quantity and distribution of neutral hydrogen is therefore intimately connected with the evolution of stars and galaxies, and observations of neutral hydrogen can give insight into these processes.  Above redshift $z=2.2$, the Ly-$\\alpha$ line redshifts into optical wavelengths and \\HI\\ can be observed, typically in absorption against distant quasars \\cite{2009ApJ...696.1543P}. Below redshift $z=0.1$, \\HI\\ has been studied using 21\\,cm emission from its hyperfine splitting \\cite{2005MNRAS.359L..30Z, 2010ApJ...723.1359M}.  There, the abundance and large-scale distribution of neutral hydrogen are inferred from large catalogs of discrete galactic emitters. Between $z=0.1$ and $z=2.2$ there are fewer constraints on neutral hydrogen, and those that do exist \\cite{2011ApJ...732...35M, 2007MNRAS.376.1357L, 2006ApJ...636..610R} have large uncertainties.  While the 21\\,cm line is too faint to observe individual galaxies in this redshift range, one can nonetheless pursue three-dimensional (3D) intensity mapping \\cite{2008PhRvL.100i1303C, 2008PhRvL.100p1301L, 2012A&A...540A.129A, 2008PhRvD..78b3529M, 2010ApJ...721..164S, 2012ApJ...752...80M}.  Instead of cataloging many individual galaxies, one can study the large-scale structure (LSS) directly by detecting the aggregate emission from many galaxies that occupy large $\\sim1000\\,{\\rm Mpc}^3$  voxels.  The use of such large voxels allows telescopes such as the Green Bank Telescope (GBT) to reach $z\\sim1$, conducting a rapid survey of a large volume.  Aside from being used to measure the hydrogen content of galaxies, intensity mapping promises to be an efficient way to study the large-scale structure of the Universe. In particular, the method could be used to measure the baryon acoustic oscillations to high accuracy and constrain dark energy \\citep{2008PhRvL.100i1303C}. However, intensity mapping is a new technique which is still being pioneered. Ongoing observational efforts such as the one presented here are essential for developing this technique as a powerful probe of cosmology.  Synchrotron foregrounds are the primary challenge to this method, because they are three orders of magnitude brighter than the 21\\,cm signal. However, the physical process of synchrotron emission is known to produce spectrally smooth radiation \\cite{2003MNRAS.346..871O, 2010ApJ...721..164S}. If the calibration, spectral response and beam width of the instrument are well-controlled and characterized, the subtraction of foregrounds should be possible because the foregrounds have fewer degrees of freedom than the cosmological signal.  We find that this allows the foregrounds to be cleaned to the level of the expected signal.  The auto-correlation of intensity maps is biased by residual foregrounds, and minimizing and constraining these residuals is an active area of work.  However, because residual foregrounds should be uncorrelated with the cosmological signal, they only boost the noise in a cross-correlation with existing surveys. This makes the cross-correlation a robust indication of neutral hydrogen density fluctuations in the 21\\,cm intensity maps \\cite{2010Natur.466..463C, 2012A&A...539L...5V}.  The first detection of the cross-correlation between LSS and 21\\,cm intensity maps at $z\\sim1$ was reported in \\citet{2010Natur.466..463C}, based on data from GBT and the DEEP2 galaxy survey.  Here we improve on these measurements by cross correlating new intensity mapping data with the WiggleZ Dark Energy Survey \\cite{2010MNRAS.401.1429D}.  Our measurement improves on the statistical precision and range of scales of the previous result, which was based on 15\\,hr of GBT integration time over 2~deg.~sq.  Throughout, we use cosmological parameters from \\citet{2009ApJS..180..330K}, in accord with \\citet{2011MNRAS.415.2876B}. ",
        "conclusions": "\\label{sec:results}  To relate the measured spectra with theory, we start with the mean 21\\,cm emission brightness temperature \\cite{2010Natur.466..463C}, \\begin{equation} T_b=0.29\\frac{\\Omega_{\\rm {HI}}}{10^{-3}}\\left(\\frac{\\Omega_m+ (1+z)^{-3}\\Omega_\\Lambda}{0.37}\\right)^{-\\frac{1}{2}}\\left(\\frac{1+z}{1.8} \\right)^{\\frac{1}{2}}~{\\rm mK}. \\end{equation} Here $\\Omega_{\\rm{HI}}$ is the comoving \\HI\\ density (in units of today's critical density), and $\\Omega_m$ and $\\Omega_\\Lambda$ are evaluated at the present epoch.  We observe the brightness contrast, $\\delta T = T_b \\delta_{\\rm{HI}}$, from fluctuations in the local \\HI\\ over-density $\\delta_{\\rm{HI}}$.  On large scales, it is assumed that neutral hydrogen and optically-selected galaxies are biased tracers of the dark matter, so that $\\delta_{\\rm{HI}} = b_{\\rm{HI}} \\delta$, and $\\delta_{\\rm{opt}} = b_{\\rm{opt}} \\delta$.  In practice, both tracers may contain a stochastic component, so we include a galaxy-\\HI\\ correlation coefficient $r$. This quantity is scale-dependent because of the $k$-dependent ratio of shot noise to large-scale structure, but should approach unity on large scales.  The cross-power spectrum is then given by $P_{\\rm{HI},\\rm{opt}}(k)=T_b b_{\\rm{HI}}b_{\\rm{opt}}rP_{\\delta\\delta}(k)$ where $P_{\\delta\\delta}(k)$ is the matter power spectrum.  The large-scale matter power spectrum is well-known from CMB measurements \\cite{2011ApJS..192...18K} and the bias of the optical galaxy population is measured to be $b_{\\rm{opt}}^2=1.48\\pm0.08$ at the central redshift of our survey \\cite{2011MNRAS.415.2876B}. Simulations including nonlinear scales (as in Sec.~\\ref{ss:datatomaps}) are run through the same pipeline as the data. We fit the unknown prefactor $\\Omega_{\\rHI}b_{\\rHI}r$ of the theory to the measured cross-powers shown in Fig.~\\ref{f:corrcomb}, and determine $\\Omega_{\\rHI}b_{\\rHI}r=[0.44\\pm0.10({\\rm stat.})\\pm0.04({\\rm sys.})]\\times 10^{-3}$ for the 15\\,hr field data, and $\\Omega_{\\rHI}b_{\\rHI}r=[0.41\\pm0.11({\\rm stat.})\\pm0.04({\\rm sys.})]\\times10^{-3}$ for the 1\\,hr field data.  The systematic term represents the $9\\%$ absolute calibration uncertainty from Sec.~\\ref{ss:datatomaps}. It does not include current uncertainties in the cosmological parameters or in the WiggleZ bias, but these are sub-dominant. Combining the two fields yields $\\Omega_{\\rHI}b_{\\rHI}r=[0.43\\pm0.07({\\rm stat.})\\pm0.04({\\rm sys.})]\\times10^{-3}$.  These fits are based on the range $0.075\\,h{\\rm Mpc}^{-1}<k<0.3\\,h{\\rm Mpc}^{-1}$ over which we believe that errors are well-estimated (failing toward larger scales where there are too few $k$ modes in the volume) and under the assumption that nonlinearities and the beam/pixelization (failing toward smaller scales) are well-understood. A less conservative approach is to fit for $0.05\\,h{\\rm Mpc}^{-1}<k<0.8\\,h{\\rm Mpc}^{-1}$ where the beam, model of nonlinearity and error estimates are less robust, but which shows the full statistical power of the measurement, at $7.4\\sigma$ combined. Here, $\\Omega_{\\rHI}b_{\\rHI}r=[0.40\\pm0.05({\\rm stat.})\\pm0.04({\\rm sys.})]\\times10^{-3}$ for the combined, $\\Omega_{\\rHI}b_{\\rHI}r=[0.46\\pm0.08]\\times10^{-3}$ for the 15\\,hr field and $\\Omega_{\\rHI}b_{\\rHI}r=[0.34\\pm0.07]\\times10^{-3}$ for the 1\\,hr field.  To compare to the result in \\citet{2010Natur.466..463C}, $\\Omega_{\\rHI}b_{\\rm rel}r=[0.55\\pm0.15({\\rm stat.})]\\times10^{-3}$, we must multiply their relative bias (between the GBT intensity map and DEEP2) by the DEEP2 bias $b=1.2$ \\citep{2004ApJ...609..525C} to obtain an expression with respect to $b_{\\rHI}$. This becomes $\\Omega_{\\rHI}b_{\\rHI}r=[0.66\\pm0.18({\\rm stat.})]\\times10^{-3}$, and is consistent with our result.  The absolute abundance and clustering of \\HI\\ are of great interest in studies of galaxy and star formation. Our measurement is an integral constraint on the \\HI\\ luminosity function, which can be directly compared to simulations. The quantity $\\Omega_{\\rHI}b_{\\rm HI}$ also determines the amplitude of 21\\,cm temperature fluctuations. This is required for forecasts of the sensitivity of future 21\\,cm intensity mapping experiments. Since $r<1$ we have put a lower limit on $\\Omega_{\\rHI}b_{\\rHI}$.  To determine $\\Omega_{\\rHI}$ alone from our cross-correlation requires external estimates of the \\HI\\ bias and stochasticity.  The linear bias of \\HI\\ is expected to be $\\sim0.65$ to $\\sim1$ at these redshifts \\cite{2010ApJ...718..972M, 2011MNRAS.415.2580K}.  Simulations to interpret \\citet{2010Natur.466..463C} find values for $r$ between $0.9$ and $0.95$ \\cite{2011MNRAS.415.2580K}, albeit for a different optical galaxy population. Measurements of the correlation coefficient between WiggleZ galaxies and the total matter field are consistent with unity in this $k$-range (with $r_{m,{\\rm opt}} \\gtrsim 0.8$) \\cite{2011MNRAS.415.2876B}.  These suggest that our cross-correlation can be interpreted as $\\Omega_{\\rHI}$ between $0.45\\times 10^{-3}$ and $0.75 \\times 10^{-3}$.  Measurements with Sloan Digital Sky Survey \\citep{2009ApJ...696.1543P} suggest that before $z=2$, $\\Omega_{\\rHI}$ may have already reached $\\sim0.4\\times10^{-3}$. At low redshift, 21\\,cm measurements give $\\Omega_{\\rHI}(z\\sim0)=(0.43\\pm0.03) \\times10^{-3}$ \\cite{2010ApJ...723.1359M}.  Intermediate redshifts are more difficult to measure, and estimates based on Mg-II lines in DLA systems observed with Hubble Space Telescope find $\\Omega_{\\rHI}(z\\sim1)\\approx(0.97\\pm0.36)\\times10^{-3}$ \\citep{2006ApJ...636..610R}, in rough agreement with $z\\approx0.2$ DLA measurements \\citep{2011ApJ...732...35M} and 21\\,cm stacking \\citep{2007MNRAS.376.1357L}.  This is in some tension with a model where $\\Omega_{\\rHI}$ falls monotonically from the era of maximum star formation rate \\citep{2012MNRAS.420.2799D}.  Under the assumption that $b_{\\rHI}=0.8, r=1$, the cross-correlation measurement here suggests $\\Omega_{\\rHI}\\sim0.5\\times 10^{-3}$, in better agreement, but clearly better measurements of $b_{\\rHI}$ and $r$ are needed.  Redshift space distortions can be exploited to break the degeneracy between $\\Omega_{\\rHI}$ and bias to measure these quantities independently of simulations \\cite{2008arXiv0804.1624W, 2010PhRvD..81j3527M}.  This will be the subject of future work.  Our measurement is limited by both the number of galaxies in the WiggleZ fields and by the noise in our radio observations. Simulations indicate that the variance observed in our radio maps after foreground subtraction is roughly consistent with the expected levels from thermal noise.  This is perhaps not surprising, our survey being relatively wide and shallow compared to an optimal LSS survey, however, this is nonetheless encouraging."
    },
    "1208/1208.2272_arXiv.txt": {
        "abstract": "{} {We present a new sample of X--ray selected galaxy groups and clusters serendipitously observed with the X--ray Telescope (XRT) on board of the Swift satellite. Using the XRT archive as of April 2010, we searched for extended sources among 336 GRB fields with galactic latitude $|b|>$20$^{\\circ}$. Our selection algorithm provides us with a flux-limited sample of 72 X--ray groups and clusters with a well defined selection function and an expected negligible contamination. The sky coverage of the survey goes from the total 40 deg$^2$ to 1 deg$^2$ at a flux limit of $10^{-14}$ erg s$^{-1}$ cm$^{-2}$ ($0.5-2$ keV). This paper provides a description of the XRT data processing, the statistical calibration of the survey, and the catalog of detected cluster candidates.} {All the X--ray sources are detected in the Swift-XRT soft ($0.5-2$ keV) images with the algorithm {\\tt wavdetect}.  A size parameter defined as the half power radius (HPR) measured inside a box of 45$\\times$45 arcsec, is assigned to each source. We select extended sources by applying a threshold on the HPR.  Thanks to extensive simulations, we are able to calibrate the threshold value, which depends on the measured net counts inside the box and on the image background, in order to identify all the sources with a probability $\\simeq 99$\\% of being extended. The net counts associated to each extended source are then computed by simple aperture photometry. } {We compute the logN--logS of our sample, finding very good agreement with previous deep cluster surveys.  We did not find any correlation between the cluster and the GRB positions. A cross correlation with published X--ray catalogs shows that only 9 sources were already detected, none of them as extended. Therefore, $\\sim 90$\\% of our sources are new X--ray detections.  We also cross correlated our sources with optical catalogs, finding 20 previously identified clusters.  Overall, about $\\sim 65$\\% of our sources are new detections, both as X--ray sources and as clusters of galaxies.} {The XRT follow--up observation of GRBs is providing an excellent serendipitous survey for groups and clusters of galaxies, mainly thanks to the low background of XRT and its constant angular resolution across the field of view.  A significant fraction of the sample ($\\sim 33$\\%) has spectroscopic or photometric redshift thanks to a cross-correlation with public optical surveys. } ",
        "introduction": "X--ray observations of clusters of galaxies over a significant range of redshifts have been used to investigate the chemical and thermodynamical evolution of the X--ray emitting Intra Cluster Medium \\citep[ICM, see][]{2004Ettori,2007Balestra,2008Maughan,2009Anderson}, and to const|rain the cosmological parameters and the spectrum of the primordial density fluctuations \\citep{2002Rosati,2005Schuecker, 2005Voit,2008Borgani,2009Vikhlinin,2010Mantz, 2011Allen}.  In this respect, X--ray surveys of clusters of galaxies represent a key tool for cosmology and the physics of large scale structure.  The need of assembling larger and larger X--ray selected cluster samples with well defined completeness criteria is one of the critical issues of present-day cosmology.  In order to build statistically complete cluster catalogs, a wide and deep coverage of the X--ray sky is mandatory.  To date, there is a remarkable lack of recent wide area X--ray surveys suitable to this scope.  Most of the existing cluster surveys are based on source samples selected by ROSAT, and confirmed through optical imaging and spectroscopic observations.  The most recent constraints on cosmological parameters from X--ray clusters are based on the \\chandra\\ follow--up of 400 deg$^2$ ROSAT serendipitous survey and of the All-Sky Survey \\citep{2009Vikhlinin,2010Mantz}.  Renewed interest in the field of cosmological tests with clusters has been recently provided by Sunyaev-Zel'dovich (SZ) surveys from the South Pole Telescope Survey \\citep{spt2012} and the Atacama Cosmology Project \\citep{act2012}.  At present, only modest improvements have been obtained with respect to constraints from WMAP7 plus baryonic acoustic oscillations plus Type Ia supernova.  The future of SZ cluster surveys is very promising, but at present an X--ray follow--up of SZ clusters is still needed, either for narrowing down the uncertainties on the cluster mass, or to firmly evaluate purity and completeness of the sample.  For example, a large effort is being invested in the X--ray follow--up with \\chandra\\ (PI B. Benson) of SZ selected clusters from the South Pole Telescope Survey. This follow--up will provide the X--ray data for 80 massive clusters spread over 2000 deg$^2$ in the redshift range 0.4$<z<$1.2.  \\begin{center} \\begin{table*} \\centering \\caption{Flux limited X--ray Cluster Surveys} \\label{surveys} \\begin{tabular}{@{}lcccc@{}}\\hline \\smallskip Name   & Flux limit       & solid angle & Number of sources & Reference \\\\ & cgs ($0.5-2$ keV)  &    deg$^2$   &                 &            \\\\ \\hline \\hline \\smallskip  SEXCLAS & $0.6 \\times 10^{-14}$  (min) & 2.1  &  19 & Kolokotronis et al. (2006)\\\\ DCS & $0.6\\times 10^{-14}$  (min) & 5.55  & 36 & Boschin (2002) \\\\ ChaMP & $1.0\\times 10^{-14}$   (min)& 13.0  & 49 & Barkhause et al. (2006) \\\\ \\bf SXCS  & \\bf $1.0\\times 10^{-14}$   (min)  & \\bf 40.0  &\\bf 72  & \\bf This work \\\\ XDCP & $1.0 \\times 10^{-14}$  (average)  & 76.0  & 22 ($z>0.9$) & Fassbender et al. (2011)  \\\\ XCLASS & $2 \\times 10^{-14}$ (min) & 90.0 & 347 & Clerc et al. (2012)\\\\ Peterson09 &  $ \\sim 0.3\\times 10^{-14}$ (min) & 163.4& 462 & Peterson et al. (2009)\\\\ XCS  &  $>300$ net cts  & 410.0  &  993  & Lloyd-Davies et al. (2011)\\\\ \\hline \\hline \\smallskip  SXDF & $0.2 \\times 10^{-14}$  (min) & 1.3  & 57 & Finoguenov et al. (2010)\\\\ COSMOS  & $0.2 \\times 10^{-14}$  (min) & 2.1 & 72  & Finoguenov et al. (2007)\\\\ XMM-BCS & $0.6 \\times 10^{-14}$   (min)     & 6.0 & 46  & Suhada et al. (2012)\\\\ XMM-LSS & $\\sim  10^{-14}$   (min)     & 11.0 & 66  & Adami et al. (2011)\\\\ \\hline \\hline \\smallskip \\end{tabular} \\tablefoot{List of X--ray flux limited cluster surveys based on \\chandra\\ or XMM data, updated at the time of writing (May 2012), plus the Swift/XRT Cluster Survey presented in this work (highlighted in bold).  Surveys based on archival data are listed in the upper part of the table, while dedicated (contiguous) surveys are shown in the lower part.  Surveys are ranked according to the total solid angle.  The total number of clusters refer to the X--ray selected only, while the quoted solid angle is the maximum covered by the survey.  Note that the limiting soft fluxes quoted in this table (when available) may refer to the lowest value in the sample (minimum) or to an average value over the entire solid angle (average).  For a full characterization of the survey depth, i.e., the sky coverage as a function of the flux, we refer to the corresponding papers.} \\end{table*} \\end{center}   New X--ray surveys of clusters of galaxies in the \\chandra\\ and XMM--Newton era are based on the compilation of serendipitous medium and deep--exposure extragalactic pointings not associated to previously known X--ray clusters.  Among these, one of the first survey was assembled by \\citet{2002Boschin}, who found 36 clusters (among them 28 new detections) in 5.55 deg$^2$ of surveyed area. Eventually, the \\chandra\\ Multiwavelength Project (ChaMP) Serendipitous Galaxy Cluster Survey \\citep{2006champs} identified about 50 cluster and group candidates from 130 archival \\chandra\\ pointings covering 13 deg$^2$. Of the 50 clusters, about 16 are expected to have redshift $z>$0.5.  More effort is devoted to the XMM-Newton archival data, also motivated by the larger solid angle and the nominal larger sensitivity.  The XMM--Newton Distant Cluster Project \\citep[XDCP,][]{2011fassbender} take advantage of an optical and IR follow--up of extended source candidate in XMM images, to identify high--$z$ clusters.  So far, the XDCP survey has yielded 22 spectroscopically confirmed clusters in the redshift range 0.9$<z\\simlt$1.6.  A key step in XDCP is the identification of high--$z$ candidates based on optical and NIR photometric technique,as well as extensive spectroscopic work. In this respect, it currently provides the largest sample of confirmed galaxy clusters at $z>$0.8, and the purity of the sample is extremely high. Another small survey (2.1 deg$^2$) conducted using XMM archival data is SEXCLAS \\citep{kolo2006}, which include 19 serendipitous detections down to $6\\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$.  The largest project based on the entire XMM archive is the XMM Clusters Survey \\citep[XCS,][]{romer01,2011lloyd-davies}, with about 1000 cluster candidates over a solid angle of 410 deg$^2$ (the largest based on current X--ray facilities).  Recently, about 500 clusters have been optically identified out of $\\sim 1000$ candidates \\citep{2012Mehrtens}.  Another survey based on the XMM archive is XCLASS, which is limited to brighter sources and aims at constraining cosmological parameters on the basis of the X--ray information only \\citep{clerc12}.  Finally, a survey combining 41.2 deg$^2$ of XMM and \\chandra\\ overlapping archival data, plus 122.2 deg$^2$ of \\chandra\\ only, has been presented in \\citet{peterson09}, for a total of 462 new serendipitous sources. Clearly, there is a significant overlap among the cluster samples derived from serendipitous XMM/\\chandra\\ surveys.   In addition, there are also ongoing dedicated, contiguous surveys. The cluster survey in the Subaru-XMM Deep Field (SXDF) reaches a depth of $2\\times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$ over 1.3 deg$^2$ with 57 X--ray clusters identified with the red-sequence technique \\citep{finog2010}.  Similar depth has been reached in the COSMOS field over 2 deg$^2$, for a total of 72 clusters \\citep{2007Finoguenov}. The XMM-Newton-Blanco Cosmology Survey project \\citep[XMM-BCS, ][]{suhada2012} is a multiwavelength X--ray, optical and mid-infrared cluster survey covered also by the South Pole Telescope and the Atacama Cosmology Telescope with the aim of studying the cluster population in a 14 deg$^2$ field.  The analysis of the first 6 deg$^2$ provided a sample of 46 clusters.  Finally, the largest contiguous survey is the XMM Large Scale Structure Survey \\citep[XMM-LSS,][]{2007XMMLSS} which is covering a region of 11 deg$^2$ with the aim of tracing the large scale structure of the Universe out to a redshift of $z\\sim$1. At present, a first data release from an area of 6 deg$^2$ consists in 66 spectroscopically confirmed clusters with 0.05$<z<$1.5 \\citep{2011Adami}.  One of the most important goal of these surveys is to find massive clusters at high $z$.  Clearly, tracing the most massive, gravitationally bound structures in the Universe up to the highest possible redshift is of paramount importance for structure formation and for cosmology.  While the redshift limit for X--ray selected clusters is $z\\sim 1.57$ (in XDCP), extended X--ray emission has been detected in optical and IR selected clusters at $z=1.75$ \\citep{2012Stanford}, $z=2.07$ \\citep{2011Gobat}, with an extreme candidate at $z\\sim 2.2$ \\citep[see ][]{2011andreon}.  The present situation is summarized in Table \\ref{surveys}.  This picture is not expected to change significantly in the next years, with a modest increase in the source statistics and in the quality of the X--ray data.  In particular, the study of distant ($z\\geq 1$) clusters relies almost exclusively on time-expensive follow--up with \\chandra.  In the case of \\chandra , the process of assembling a wide and deep survey is slow due to the small field of view (FOV) and the low collecting area.  In the case of XMM-Newton the collecting area is significantly higher, but the identification of extended sources, in particular at medium and high redshift, is more difficult due to the larger size of the Point Spread Function (PSF, whose half energy width is 15\" at the aimpoint) and its degradation as a function of the off--axis angle.  This is not an issue for nearby or medium redshift clusters and groups, but it becomes a problem at $z>1$, where the optical and IR data play a dominat role in cluster identifications. In addition, the relatively high and unstable background may hamper the proper characterization of low surface brightness sources.  In conclusion, while their design is optimized to obtain detailed images of isolated sources to explore the deep X--ray sky, a substantially different mission strategy is required for surveys.  In this work we present a new X--ray cluster survey using the rich archive of the X--ray Telescope (XRT) onboard of the Swift satellite \\citep{2005bur} which has never been used for the detection of extended sources.  Despite its low collecting area (about one fifth of that of \\chandra\\ at 1 keV, see Figure \\ref{xrt_arf}), XRT has two characteristics which are optimal for X--ray cluster surveys: a low background and an almost constant PSF across the FOV \\citep{2007Moretti}.  Moreover, almost all the Swift/XRT pointings can be used to build a serendipitous survey.  In fact, the operation mode of the XRT observations considered in this work, consists of a prompt follow--up of fields centered on Gamma Ray Bursts (GRB) detected by Swift.  As we will show in \\S~\\ref{cluster_cat}, the GRBs do not show any correlation with our extended sources.  Therefore, despite its small size, the XRT can be successfully used to build an unbiased cluster survey.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{xrt_arf.eps} \\caption{Effective area of XRT (green solid line) as a function of the energy compared with that of the \\chandra\\ satellite (red dotted and cyan dashed lines for \\chandra\\ ACIS-S and ACIS-I, respectively).} \\label{xrt_arf} \\end{figure}  The catalog of extended sources presented in this paper constitutes the Swift X--ray Cluster Survey (SXCS) and it is based on the 336 GRB fields with galaxtic latitude $|b|>$20$^\\circ$ present in the XRT archive as of April 2010.  The sources are identified thanks to a very simple but effective algorithm to select extended X--ray sources in the XRT soft band images.  In this paper we adopt a conservative detection threshold which guarantees very low contamination and a well defined completeness function.  This catalog is complementary to the catalog of point sources identified in the GRB fields of XRT by \\citet{2011Puccetti}, which focus on the study of the AGN population.  The paper is organized as follows. In \\S~\\ref{sur_sec} we provide a description of the principal characteristics of the XRT. In \\S~\\ref{field_sel} we describe the field selection and the data reduction.  In \\S~\\ref{ext_sel} we describe the detection algorithm and the selection of the extended source candidates.  In \\S~\\ref{ext_phot} we describe how we performed extended source photometry. In \\S~\\ref{cluster_cat} we present the final list of groups and clusters in our sample, compute the sky coverage and the LogN--LogS, and check for possible selection bias.  In \\S~\\ref{cross_corr} we correlate our catalog with existing databases to identify previously known sources and collect the spectroscopic or photometric redshifts of member galaxies available in the literature. In \\S~\\ref{discus} we discuss our results in the context of current and future X--ray surveys.  Finally, in \\S~\\ref{conc} we summarize our conclusions. ",
        "conclusions": "\\label{conc}  We present a new sample of X--ray selected groups and clusters of galaxies obtained with the X--ray Telescope (XRT) on board of the Swift satellite.  We search for extended sources among 336 GRB fields imaged with XRT with galactic latitude $|b|>$20$^\\circ$ (available in the archive as of April 2010). We identify extended sources with a simple criterion based on the measurement of the HPR within a box of $45\\times 45$ arcsec of the source image.  We apply a sharp threshold of 100 net soft counts within the extraction radius $R_{ext}$ to select reliable extended sources.  Extensive simulations showed that our method, despite being very simple, provide us with an X--ray sample with an expected high completeness and a low contamination, also thanks to a careful {\\sl a posteriori} visual inspection.  Our final group and cluster catalog consists of 72 X--ray sources. The sky coverage of the survey goes from the total 40 deg$^2$ to 1 deg$^2$ at a flux limit of about $10^{-14}$ erg s$^{-1}$ cm$^{-2}$. The corresponding logN--logS is in very good agreement with previous deep surveys.  We directly verified that there is no correlation between the position of the clusters and that of the GRBs. In other words, selecting our cluster sample from XRT GRB fields does not result in any spatial bias with respect to the GRB positions.  A cross correlation with X--ray catalogs shows that only 9 SXCS sources were previously identified in the X--ray band, none of them classified as extended.  A search in optical databases (mostly based on SDSS data) allows us to find the counterparts of 20 clusters.  In addition, 4 galaxies with redshift have been found to be within 7 arcsec from the X--ray centroid, and therefore they are considered as possible identification of the central galaxy of the group/cluster candidate.  Overall, only 20 sources are confirmed as clusters using data from the literature, while 30 sources have some counterpart in the NASA Extragalactic Database, and, finally, 42 sources are newly identified.  We estimate that about one third of the sample is detected with a S/N high enough to allow the measure of the redshift from the X--ray spectral analysis \\citep[see][]{2011Yu}, as we will show in a companion paper (Moretti et al. in preparation).  Thanks to the quality of our X--ray data, and thanks to the synergy with other surveys, specifically the SDSS, we are able to provide a X--ray cluster sample with a well established completeness function down to a flux limit comparable to that of the deepest cluster surveys based on ROSAT data \\citep[RDCS, see][]{1998Rosati}, and with a comparable statistics.  In addition, we expect to detect a few clusters with redshift $z\\geq$1.  Overall, the SXCS is expected to give a significant contribution in the field of X--ray clusters surveys also thanks to its peculiar properties of a low background and a constant PSF.  These properties are two key requirements for  future wide--area, X--ray surveys, as foreseen by proposed future missions aiming at bringing the X--ray sky to the same depth and richness of the optical and IR sky in the next decade. A deeper, extended release of the SXCS based on a new detection algorithm tailored to XRT images is currently undergoing (Liu et al. in preparation).  Catalog and data products of SXCS, constantly updated, are made avalilable to the public through the website {\\tt http://adlibitum.oats.inaf.it/sxcs}."
    },
    "1208/1208.5178_arXiv.txt": {
        "abstract": "{The  Milky Way  (MW) bulge  is a  fundamental Galactic  component for understanding the formation and evolution of galaxies, in particular our  own.  The  ESO Public  Survey  VISTA Variables  in the  V\\'{i}a L\\'actea is  a deep  near-IR survey mapping  the Galactic  bulge and southern plane.   Particularly for the  bulge area, VVV  is covering $\\sim  315$~deg$^2$.  Data taken  during 2010  and 2011  covered the entire bulge area in the $JHK_{\\rm s}$ bands.} {We used VVV data for the whole bulge area as a single and homogeneous data  set to  build for  the  first time  a single  colour-magnitude diagram (CMD) for the entire Galactic bulge.} {Photometric data in the $JHK_{\\rm  s}$ bands were combined to produce a single and  huge data set containing $173,150,467$  sources in the three  bands, for  the  $\\sim  315$~deg$^2$ covered  by  VVV in  the bulge. Selecting only the data  points flagged as stellar, the total number of sources is $84,095,284$.} {We built the largest  colour-magnitude diagrams published up to date, containing 173.1+ million sources for all data points, and more than 84.0 million  sources accounting for  the stellar sources  only. The CMD  has a  complex shape,  mostly owing  to the  complexity  of the stellar  population  and the  effects  of  extinction and  reddening towards  the Galactic  centre. The  red clump  (RC) giants  are seen double in magnitude at  $b\\sim-8^\\circ-10^\\circ$, while in the inner part  ($b\\sim-3^\\circ$) they appear  to be  spreading in  colour, or even splitting into a secondary peak. Stellar population models show the predominance of main-sequence  and giant stars.  The analysis of the outermost bulge  area reveals a well-defined sequence  of late K and  M dwarfs,  seen at  $(J-K_{\\rm s})\\sim0.7-0.9$~mag  and $K_{\\rm s}\\gtrsim14$~mag.} {The interpretation of the  CMD yields important information about the MW bulge, showing  the fingerprint of its structure  and content. We report  a well-defined red  dwarf sequence  in the  outermost bulge, which is important  for the planetary transit searches  of VVV.  The double RC in  magnitude seen in the outer bulge  is the signature of the X-shaped MW bulge, while the  spreading of the RC in colour, and even its  splitting into a  secondary peak, are caused  by reddening effects.   The  region  around  the  Galactic centre  is  harder  to interpret  because   it  is  strongly  affected   by  reddening  and extinction.} ",
        "introduction": "The bulge  of the Milky Way  (MW) is a  fundamental Galactic component for understanding the formation and  evolution not only of our Galaxy, but of galaxies  in general.  In the MW the  stars can be individually resolved down  to faint magnitudes, allowing  a complete understanding of its structure, age, kinematics, and chemical composition.  However, because the bulge and plane concentrate  not only most of the stars in the MW, but also gas and dust, observations of the Galactic centre are difficult  because  they  are  affected by  crowding  and  extinction. Optical  surveys such  as MACHO  \\citep{1999ApJS..124..171A}  and OGLE \\citep{1993AcA....43...69U}  are highly  affected by  extinction while past  near-infrared (near-IR)  surveys such  as 2MASS  are  limited to bright             sources            \\citep[$K_{\\rm            s}\\sim 14.3$~mag;][]{2006AJ....131.1163S}, and do not allow a complete view of the bulge populations.  \\begin{figure*} \\includegraphics[bb=5cm -3.0cm 16cm 10cm,angle=-90,scale=0.55]{fig01.eps} \\caption{Left-hand panels: total number of sources present in the $J$, $H$, and $K_{\\rm s}$ singleband  and in the $JHK_{\\rm s}$ multi-band catalogue for each VVV bulge tile.  The tiles are numbered from b201 to b396,  starting at the bottom-left  corner of the  VVV bulge area ($l,b=-10,-10$).   The Galactic  latitude  for the  tile centres  is shown in the  figure, while dashed lines mark the  edge of each line of  tiles  across  the  bulge.   The  right-hand  panel  presents  a histogram  with the  ratio  between  the number  of  sources in  the multi-band catalogue  in relation to its  least numerous single-band catalogue.} \\label{fig:stat} \\end{figure*}   The  new ESO  Public Survey  VISTA Variables  in the  V\\'{i}a L\\'actea (VVV)  is a  deep near-IR  survey mapping  562 square  degrees  in the southern      plane       and      bulge      of       our      Galaxy \\citep{2010NewA...15..433M}. Particularly  for the bulge  area, VVV is covering $\\sim  315$~deg$^2$. The VVV survey observes  in five near-IR passbands ($ZYJHK_{\\rm  s}$), where extinction effects  are lower than in  the optical  wavelengths,  and  it is  much  deeper than  previous near-IR  surveys.  The 5$\\sigma$  limiting magnitude  of the  VVV data using aperture photometry is  $K_{\\rm s}\\sim18.5$~mag in clean fields, which allows one to monitor bright  sources such as RR Lyrae and clump giants  stars along  the whole  Galactic bulge  and  plane, solar-type stars at the  Galactic centre distances, and even  faint dwarf stars a few kpc  away. The  VVV data reach  the main-sequence  turn-off (MSTO) even      in      intermediate-      to     high-extinction      areas \\citep[$E(B-V)\\lesssim4.0$,][]{2010NewA...15..433M}.  The  VVV photometry  for  the whole  bulge  area was  combined into  a single, huge, and  homogeneous data set, allowing us  to build for the first time  a single colour-magnitude diagram for  the entire Galactic bulge ($\\sim  315$~deg$^2$).  Here we present the  84-million star VVV colour-magnitude diagram  (CMD) for the  Galactic bulge.  This  is the largest  CMD  ever published  for  a  large  homogeneous data  set,  a significant step  forward since the publication of  the 9-million star CMD  of the  Large Magellanic  Cloud  \\citep{2000AJ....119.2194A}.  We discuss  the  differences  in   the  morphology  caused  by  crowding, extinction  and sky  brightness. A  forthcoming paper  will  present a similar  analysis for the  populations across  the VVV  Galactic plane area.  \\begin{figure*} \\includegraphics[bb=1cm -1.6cm 20cm 12cm,angle=-90,scale=0.60]{fig02.eps} \\caption{Density  plot  in logarithmic  scale  showing  the VVV  bulge area. The map was made  using the multi-band $JHK_{\\rm s}$ CASU v1.1 and  v1.2  catalogues  for   point  sources  brighter  than  $K_{\\rm s}=16.5$~mag (see  text).  Crowded  areas appear in  yellow, while less populated regions  as well as highly extincted  areas are shown in blue.  The overlapping  regions between the tiles are highlighted since the point  sources are accounted for twice.  Dashed lines mark the 15  areas selected to build the  colour-magnitude diagrams shown in Fig.~\\ref{all_15} (see Section  4).  The source density, in units of 10$^5$~sources~deg$^{-2}$, is indicated  in the horizontal bar at the top. }. \\label{density} \\end{figure*} ",
        "conclusions": "We presented  the VVV colour-magnitude diagram of  the Galactic bulge, the largest CMD ever published  for a large homogeneous data set, with 84M+  stars.    The  interpretation   of  the  CMD   yields  important information  about  the  MW  bulge,  showing the  fingerprint  of  its structure and content.  Stellar  population synthesis models  fit the  data well,  showing the predominance  of main-sequence  and giant  stars in  the  outer bulge, which belong both  to the bulge and halo.   Thin- and thick-disk stars are also  present in  fewer numbers.   In the inner  bulge the  CMD is dominated by  bulge giants, which contribute  up to 46\\%  of the total sources at $b\\sim-3^\\circ$  The  analysis  of the  outermost  bulge  area  reveals a  well-defined sequence   of   late   K    and   M   dwarfs,   seen   at   $(J-K_{\\rm s})\\sim0.7-0.9$~mag  and $K_{\\rm  s}\\gtrsim14$~mag. These  stars are particularly important  in the  planetary transit searches  during the variability campaign of VVV.  The RC appears to be double in magnitude in the outer bulge region due to the X-shaped structure of the MW bulge, with the separation between the peaks  reaching $\\Delta  K_{\\rm s}=0.70$ at  $b\\sim-9^\\circ$. This result  is  complementary with  previous  analyses,  and confirms  the persistence  of the  X-shaped structure  for  $b\\lesssim-8^\\circ$.  In contrast, in  the inner  part ($b\\sim-3^\\circ$) the  RC appears  to be spreading in  colour, or even splitting  into a secondary  peak with a separation of $\\Delta(J-K_{\\rm s})\\gtrsim~0.2$~mag (area \\#9). The use of  reddening-free  parameters confirmed  that  this  last feature  is caused by reddening effects.  The CMDs of  the Galactic centre are harder  to interpret because they are strongly affected by reddening and extinction. All populations are seen to be much redder, and  the RC spreads along the direction of the reddening vector by $\\Delta(J-K_{\\rm s})\\gtrsim~2.0$~mag."
    },
    "1208/1208.5949_arXiv.txt": {
        "abstract": "We present nine near-infrared (NIR) spectra of supernova (SN) 2005cf at epochs from $-10$\\,d to $+42$\\,d with respect to $B$-band maximum, complementing the existing excellent data sets available for this prototypical Type Ia SN at other wavelengths. The spectra show a time evolution and spectral features characteristic of normal Type Ia SNe, as illustrated by a comparison with SNe~1999ee, 2002bo and 2003du. The broad-band spectral energy distribution (SED) of SN~2005cf is studied in combined ultraviolet (UV), optical and NIR spectra at five epochs between $\\sim$\\,8\\,d before and $\\sim$\\,10\\,d after maximum light. We also present synthetic spectra of the hydrodynamic explosion model W7, which reproduce the key properties of SN~2005cf not only at UV-optical as previously reported, but also at NIR wavelengths. From the radiative-transfer calculations we infer that fluorescence is the driving mechanism that shapes the SED of SNe~Ia. In particular, the NIR part of the spectrum is almost devoid of absorption features, and instead dominated by fluorescent emission of both iron-group material and intermediate-mass elements at pre-maximum epochs, and pure iron-group material after maximum light. A single P-Cygni feature of \\MgII\\ at early epochs and a series of relatively unblended \\CoII\\ lines at late phases allow us to constrain the regions of the ejecta in which the respective elements are abundant. ",
        "introduction": "An empirical relation between their maximum brightness and light curve decline rate \\citep{phillips1993a} makes Type Ia supernovae (SNe~Ia) one of the best tools to measure cosmological distances. Using this approach \\citet{riess1998a} and \\citet{perlmutter1999a} discovered that our Universe undergoes an accelerated expansion. This was in turn interpreted as evidence for the existence of `dark energy' (e.g.~\\citealt{riess2007a}). Despite this importance for astrophysics and cosmology our knowledge about the progenitor systems or the explosion mechanisms of these luminous events is still limited. They are commonly believed to be thermonuclear explosions of carbon-oxygen white dwarfs in binary systems. However, many aspects are still unclear; e.g. the masses of the exploding stars, the nature of the companion star, the ignition conditions, the exact mechanism of the burning process or the origin of the various subclasses of SNe~Ia. Systematic uncertainties arising from this ignorance are one factor limiting the precision of cosmological measurements. One way to improve our understanding of these objects is the comparison of observed spectra with predictions from theoretical models. This, however, requires complete data sets with a wide wavelength coverage \\citep{roepke2012a}.  In particular, the near-infrared\\footnote{In this paper we will refer to wavelengths between 9500 and 26\\,000\\,\\AA\\ when talking about `near-infrared'.} (NIR) regime has been proposed as a useful tool to discriminate between different explosion models \\citep[e.g.][]{marion2003a}. This is mainly due to the fact that the lower opacity in the NIR enables us to look deeper into the ejecta and thus probe layers still hidden to optical observations at the same epochs. Also, the NIR is believed to be less severely affected by line blending compared to optical or ultraviolet (UV) wavelengths \\citep{harkness1991,hoeflich1993a,spyromilio1994,wheeler1998}. At the same time, however, observations suggest SNe~Ia to be excellent NIR standard candles (\\citealt{elias1985a,meikle2000, krisciunas2004a}, but see \\citealt{kattner2012a}), in contrast to optical wavelengths where they can only be used as standard\\textit{izeable} candles. This result was confirmed by a series of toy explosion models \\citep{kasen2006b}. While the homogeneity in the NIR might point towards a common explosion mechanism for all SNe~Ia, it could also mean that the NIR photons are not as sensitive to variations in the ejecta as previously suggested.   While optical spectra have been studied quite extensively, the NIR regime is still comparatively unexplored. Although there are quite a few publications containing NIR spectra of SNe~Ia (e.g. \\citealt{frogel1987a,spyromilio1992a,spyromilio2004a,meikle1996a, bowers1997,hamuy2002a,hamuy2002b,hoeflich2002b,rudy2002,marion2003a, marion2009,benetti2004a,eliasrosa2006a,stanishev2007b,taubenberger2011}) one has to realize that for many SNe there is only one spectrum available -- typically at post-maximum epochs -- or that the objects are of a peculiar subclass. The number of well-sampled \\textit{normal} SNe~Ia with pre-maximum-light observations in the NIR is actually quite small. \\citet{meikle1996a} have published six spectra of SN~1994D, \\citet{hamuy2002a,hamuy2002b} eleven spectra of SN~1999ee, \\citet{benetti2004a} six spectra of SN~2002bo, \\citet{pignata2008} four spectra of SN~2002dj, \\citet{eliasrosa2006a} ten spectra of SN~2003cg and \\citet{marion2009} four spectra of SN~2005am. Moreover, there are 14 spectra of SN~2003du out to the nebular phase, published by \\citet{motohara2006}, \\citet{stanishev2007b} and \\citet{marion2009}.  In order to increase the sample of SNe~Ia with good NIR coverage we present and analyze a set of nine NIR spectra of SN~2005cf, ranging from $-10.2$\\,d to $+41.5$\\,d with respect to $B$-band maximum. SN~2005cf, according to \\citet{wang2009} `the golden standard Type Ia supernova', is a spectroscopically normal SN~Ia with a more complete data set than most other SNe~Ia. Optical and NIR photometry as well as optical spectroscopy of SN~2005cf have been published by \\citet{garavini2007}, \\citet{pastorello2007a} and \\citet{wang2009}. UV photometry and spectroscopy have been published by \\citet{bufano2009} and \\citet{wang2012a}.  Table \\ref{table:SN2005cf_properties} summarizes the most important properties of SN~2005cf: at $B$-band maximum, it had a magnitude of $M_{B_{\\mathrm{max}}} = -19.39$ \\citep{pastorello2007a}, $\\Delta \\mathrm{m}_{15}(B)$ lies between 1.07 \\citep{wang2009} and 1.12 \\citep{pastorello2007a}, and the derived \\textsuperscript{56}Ni mass is about 0.7\\,$\\mathrm{M}_{\\astrosun}$ \\citep{pastorello2007a} or 0.78\\,$\\mathrm{M}_{\\astrosun}$ \\citep{wang2009}. This corresponds quite well to the average values obtained for normal SNe~Ia. An important shortcoming of the data set of SN~2005cf has so far been the lack of published NIR spectroscopy. This deficiency is now remedied with the data presented in this work.  \\begin{table} \\caption{Properties of SN~2005cf, as derived in the literature.} \\centering \\begin{tabular}{@{}lcc@{}} \\hline & \\citeauthor{pastorello2007a} & \\citeauthor{wang2009} \\\\ & \\citeyearpar{pastorello2007a} & \\citeyearpar{wang2009} \\\\ \\hline JD of $B_{\\mathrm{max}}$ & 2\\,453\\,534.0 $\\pm$ 0.3& 2\\,453\\,533.66 $\\pm$ 0.28 \\\\ $m_{B_{\\mathrm{max}}}$ & 13.54 $\\pm$ 0.02 & 13.63 $\\pm$ 0.02 \\\\ $M_{B_{\\mathrm{max}}}$ & $-19.39$ $\\pm$ 0.33 &  \\\\ $\\Delta \\mathrm{m}_{15}(B)_{\\mathrm{true}}$ & 1.12 $\\pm$ 0.03 & 1.07 $\\pm$ 0.03 \\\\ $M(^{56}\\mathrm{Ni})$ & 0.7\\,$\\mathrm{M}_{\\astrosun}$ & (0.78 $\\pm$ 0.10)\\,$\\mathrm{M}_{\\astrosun}$ \\\\ $\\mu$(MCG-01-39-003) & (32.51 $\\pm$ 0.33)\\,mag &  (32.31 $\\pm$ 0.11)\\,mag \\\\ $E(B - V)_{\\mathrm{host}}$ & 0\\,mag & (0.10 $\\pm$ 0.03)\\,mag \\\\ \\hline \\end{tabular} \\label{table:SN2005cf_properties}  \\end{table}  The paper is organised as follows. In Section~\\ref{Observations} the NIR spectra of SN~2005cf are presented and the process of data reduction is described. In Section~\\ref{Discussion} we compare these observations to the NIR spectra of other SNe~Ia and construct combined UV-optical-NIR spectra to investigate the broad-band spectral energy distribution (SED). Synthetic spectra of the hydrodynamic explosion model W7 \\citep{nomoto1984a} are used to explain the SED in terms of fluorescence and to perform a NIR line identification (Section~\\ref{Comparison_SN2005cf_W7}). Conclusions are drawn in Section~\\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}  We have presented NIR spectroscopy of SN~2005cf covering epochs from $-10$\\,d to $+42$\\,d with respect to $B$-Band maximum. Together with archival data at other wavelengths this makes SN~2005cf one of the currently best-observed SNe~Ia and allowed us to compile combined UV-optical-NIR spectra at five epochs between $\\sim$\\,8\\,d before and $\\sim$\\,10\\,d after maximum light. In this respect SN~2005cf is a unique object since to our best knowledge thus far only a single UV-through-NIR spectrum of a SN~Ia has been published by \\citet{foley2012b} for SN~2011iv.  A comparison of our spectra to observations of SNe~1999ee \\citep{hamuy2002a,hamuy2002b}, 2002bo \\citep{benetti2004a} and 2003du \\citep{stanishev2007b} shows that SN~2005cf is a perfectly normal SN~Ia at NIR wavelengths. A similar conclusion has already been drawn at optical wavelengths by \\citet{garavini2007}, \\citet{pastorello2007a} and \\citet{wang2009}.  We have also performed radiative-transfer simulations to obtain NIR spectra for the standard hydrodynamical explosion model W7 \\citep{nomoto1984a}. Comparing the obtained synthetic spectra to our spectra of SN~2005cf we find that W7 reproduces the key properties of normal SNe~Ia not only at UV-optical as reported in earlier work \\citep[e.g.][]{branch1985a,kasen2006a,kromer2009a,foley2012b} but also at NIR wavelengths. From a detailed analysis of the emission processes in the radiative-transfer simulation we have identified the main spectral features in the NIR. We find only a single clear P-Cygni line: the \\MgII\\ $\\lambda$10926 triplet which is present at early epochs. Apart from that the NIR part of the spectrum is almost devoid of absorption features. Instead, fluorescence is the driving mechanism that shapes the SED of SNe~Ia.  Before maximum light both intermediate-mass elements and iron-group elements contribute to the fluorescent emission. After maximum light almost all the flux originates from iron-group elements, though different iron-group species contribute to the various NIR bands. This makes the $J-H$ and $J-K$ colours a few weeks after maximum light potential probes to the ratio of different iron-group species. While most of the fluorescent emission is strongly blended, we find that the characteristic \\CoII\\ humps which form at about a month after maximum light in the $K$ band result from a series of relatively unblended lines. This provides a simple possibility to constrain the iron-group-element-rich regions of the ejecta from observed spectra."
    },
    "1208/1208.1667_arXiv.txt": {
        "abstract": "{Bipolar magnetic regions are formed when loops of magnetic flux emerge at the solar photosphere. Magnetic buoyancy plays a crucial role in this flux emergence process, particularly at larger scales. However it is not yet clear to what extent the local convective motions influence the evolution of rising loops of magnetic flux.} {Our aim is to investigate the flux emergence process in a simulation of granular convection. In particular we aim to determine the circumstances under which magnetic buoyancy enhances the flux emergence rate (which is otherwise driven solely by the convective upflows).} {We use three-dimensional numerical simulations, solving the equations of compressible magnetohydrodynamics in a horizontally-periodic Cartesian domain. A horizontal magnetic flux tube is inserted into fully developed hydrodynamic convection. We systematically vary the initial field strength, the tube thickness, the initial entropy distribution along the tube axis and the magnetic Reynolds number.} {Focusing upon the low magnetic Prandtl number regime ($Pm<1$) at moderate magnetic Reynolds number, we find that the flux tube is always susceptible to convective disruption to some extent. However, stronger flux tubes tend to maintain their structure more effectively than weaker ones. Magnetic buoyancy does enhance the flux emergence rates in the strongest initial field cases, and this enhancement becomes more pronounced when we increase the width of the flux tube. This is also the case at higher magnetic Reynolds numbers, although the flux emergence rates are generally lower in these less dissipative simulations because the convective disruption of the flux tube is much more effective in these cases. These simulations seem to be relatively insensitive to the precise choice of initial conditions: for a given flow, the evolution of the flux tube is determined primarily by the initial magnetic field distribution and the magnetic Reynolds number.} {} ",
        "introduction": "It is generally believed that most of the large-scale toroidal magnetic field in the solar interior is stored in the stably-stratified layer just below the base of the convection zone, where the toroidal field is amplified by the effects of differential rotation \\citep[see, for example,][]{OSSEN03}. Through the action of magnetic buoyancy \\citep[][]{PARK55}, segments of this toroidal flux rise into the convection zone, eventually emerging to form bipolar active regions at the solar surface \\citep[][]{ZWAAN85}. Several numerical simulations of this flux emergence process have focused upon the interactions between convection and buoyant magnetic flux tubes \\citep[][]{DORCH01,FAN03,ABBETT04,STEIN11}, whilst  others have focused upon the subsequent small-scale and large-scale evolution of the emerging field, as it  rises into the photosphere/chromosphere \\citep[e.g.][]{CSM07, TMI09} and into the lower corona \\citep[e.g.][]{MAG01,FAN01,ARC04,SYK08}. There have also been a number of recent reviews of flux emergence dynamics \\citep[e.g.][]{HOOD11,ARC12}.  \\par Since it is not yet possible to model the whole flux emergence process (from the base of the solar convection zone into the solar atmosphere) in a self-consistent way, it is necessary to consider local models with idealised initial conditions. Usually, the initial magnetic field distribution corresponds to a thin horizontal flux tube. This flux tube is often twisted, although the components of the magnetic field that are perpendicular to the tube axis are usually much weaker than the axial component. As a result, the dominant component of the Lorentz force corresponds to the radial gradient of the magnetic pressure. In calculations of this type, the local gas pressure is usually modified so that this initial flux tube is in approximate (total) pressure balance with its non-magnetic surroundings. The tube is then made buoyant by specifying a particular entropy distribution (or, equivalently, a density deficit) along its axis. In some models, the imposed specific entropy distribution varies along the axis of the tube, with the result that only one part of the tube is buoyant. This differential buoyancy  leads to the development of $\\Omega$-like loops of magnetic flux. When convective motions are present, it is not necessary to impose a variable entropy distribution along the tube in order to obtain $\\Omega$-like loops. Convective upflows (and downflows) naturally lead to the formation of loop-like structures, provided that the magnetic Reynolds number is large enough to ensure that the magnetic field lines are (at least partially) advected with the flow.  \\par In the absence of convective motions, there is nothing to disrupt the buoyant rise of a magnetic flux tube. The situation is clearly more complicated when the tube rises through a convective layer. In a previous study, \\citet{FAN03} used the anelastic approximation to model the interaction between uniformly buoyant magnetic flux tubes and convection in the deep layers of the solar convection zone. Varying the initial field strength in their model, they found that weak fields tended to be disrupted by the convective motions. In fact the peak field strength within the tube had to be significantly greater than the equipartition value (at which the field is in energy balance with the surrounding motions) before magnetic buoyancy could play a dominant role in the evolution of the flux tube. Similar conclusions were reached by \\citet{ABBETT04} in their related study. Although we do not go into more details here, it is worth noting one further aspect of these anelastic models. Defining the magnetic Prandtl number, $Pm$, to be the ratio of the magnetic Reynolds number to the (fluid) Reynolds number, \\citet{FAN03} focused exclusively upon the $Pm>1$ regime, at relatively low Reynolds numbers. The simulations of \\citet{ABBETT04} were at higher Reynolds number, but most were still in this $Pm>1$ regime. In actual fact, we would expect $Pm\\ll 1$ throughout the solar convection zone \\citep[see, e.g.][]{OSSEN03}. However, the low magnetic Prandtl number regime is numerically extremely challenging, which presumably explains why these studies did not focus upon this region of parameter space.  \\par Models of flux emergence through fully compressible convection have also been carried out. It is possible to simulate this process on scales that are comparable to that of an active region \\citep[][]{CRTS10,STEIN11}. However, it is not feasible to carry out parametric surveys in calculations of this size, so most previous studies have focused upon flux emergence through convective layers with a relatively small number of granular cells. The key finding of the compressible model of \\citet{DORCH01} is that convective disruption removes magnetic flux from the tube as it rises, thus reducing the quantity of flux that emerges at the surface layers. In some sense, therefore, convective motions reduce the efficiency of the flux emergence process. However, it is not clear to what extent this effect is parameter-dependent. \\citet{CSM07} considered a model of compressible convection that included the effects of radiative transfer. At some instant in time, they introduced a long, uniformly-buoyant magnetic flux tube into the lower part of the domain, modifying the local gas pressure and the local velocity field so that the tube was (at least instantaneously) in equilibrium with its non-magnetic surroundings. Under the action of the vertical convective motions, they found that weak flux tubes tend to develop a sea-serpent-like form. However, as in the anelastic studies, magnetic buoyancy dominates the evolution of the tube only when the peak magnetic field strength exceeds some threshold value. At the surface these emerging fields can be strong enough to modify the granulation pattern. In their larger scale simulations, \\citet{STEIN11} adopted a different approach, introducing untwisted, horizontal magnetic flux through the lower boundary of their convective domain, ensuring that it had the same entropy as the convective upflows. Like \\citet{CSM07}, they found that weak fields were susceptible to convective disruption. Fields that were strong enough to rise to the surface through the action of magnetic buoyancy tended to modify the surrounding flow, producing unrealistically hot, large granules.  \\par Motivated by these previous studies, the aim of this paper is to investigate the parametric dependence of simulations of flux emergence across a convectively-unstable compressible layer in the ``low\" magnetic Prandtl number regime. Most of the compressible calculations that are described above utilise artificial viscosities to stabilise the numerical scheme, so it is difficult to define a magnetic Prandtl number in these cases although, in practice, it is likely that $Pm \\approx 1$ in these calculations. Therefore, this will be the first time that the $Pm <1$ regime has been studied systematically in this context. Given that {\\it any} initial conditions in local models of this type are likely to be highly idealised, it is important to determine the extent to which the evolution of the system is sensitive to the precise choices that are made. We therefore consider several different initial configurations, varying the peak field strength, the width (and twist) of the tube, in addition to varying the initial entropy distribution within the flux tube. By carrying out this systematic survey we hope to identify the key features of the initial configuration that determine whether or not magnetic buoyancy contributes towards the evolution of the tube. The paper is structured as follows. In the next section, we describe the governing equations, the model parameters and the choice of initial conditions under consideration. In Section 3, we present our numerical results. In the last section, we summarise our findings and discuss some of their implications for flux emergence in the solar convection zone. ",
        "conclusions": "We have used numerical simulations to investigate the dynamical evolution of a magnetic flux tube embedded within a granular convective layer. Throughout the calculations we have adopted the same basic hydrodynamic flow, with a Reynolds number of $\\mathcal{R}e \\approx 420$. When the magnetic flux tube is introduced, the gas pressure is adjusted so that the tube is in pressure balance with its surroundings. The tube is then made buoyant by choosing an appropriate entropy distribution. The subsequent evolution of this tube is influenced by two key processes, namely convective disruption and magnetic buoyancy. However the extent to which magnetic buoyancy plays a role in the flux emergence process depends in a rather subtle way upon the model parameters. In this systematic survey, we varied the initial peak magnetic field strength, the twist and width of the flux tube, the magnitude of the entropy along the tube axis and the magnetic Reynolds number. In all cases, the magnetic Prandtl number, $Pm=\\mathcal{R}m/\\mathcal{R}e$ is less than unity, as would be expected for the solar convection zone. This is the first study of this type to explore this low $Pm$ regime.  \\par At a moderate magnetic Reynolds number of $\\mathcal{R}m\\approx 140$, the evolution of a thin magnetic flux tube is partially determined by the effects of convective disruption. In the weak field (high $\\beta$) cases, the magnetic field appears to play a relatively passive role in the dynamics. However, there is an enhanced flux emergence rate in the strongest field cases due to the effects of magnetic buoyancy. Increasing the magnitude of the initial peak entropy along the tube axis does not change the $\\beta$ dependence, but it does increase the rate of flux emergence in all cases. Varying the twist of the flux tube does not seem to influence the evolution of the system (presumably because the magnetic tension is still much weaker than the magnetic pressure gradients in the flux tube). However, increasing the width of the tube does produce an enhanced flux emergence rate in the strong field case. This is presumably because a thicker tube is more robust, and tends to resist convective disruption more effectively than a thinner flux tube. As a result of this, the flux tube maintains its coherence more effectively during the early stages of evolution, which allows magnetic buoyancy to operate in a more efficient manner. So our main conclusion from the $\\mathcal{R}m\\approx 140$ cases is that the contribution of magnetic buoyancy to the flux emergence process is directly related not only to the peak field strength, but also to the width of the initial magnetic flux tube. \\par We also varied the magnetic Reynolds number in this system in order to determine the extent to which the evolution of the tube is influenced by the effects of magnetic dissipation. This parametric survey produced two key results. Firstly, magnetic buoyancy contributes most effectively to the flux emergence process in the low $\\beta$, high magnetic Reynolds number regime. Secondly, for a given value of $\\beta$, the flux emergence rate is always lower in the higher $\\mathcal{R}m$ regime. This is due to the fact that convective disruption is much more efficient at higher magnetic Reynolds number (where field lines are more easily advected by the flow). The extent to which this result depends upon the magnetic Prandtl number is unclear, although moving to higher $\\mathcal{R}m$ (thereby increasing $Pm$) would be computationally simpler than reducing $Pm$ by increasing $\\mathcal{R}e$. At higher $\\mathcal{R}m$ this convective flow should be capable of sustaining a small-scale dynamo, although it is not clear how such a flow would interact with an emerging flux tube. This is a possible area for future work.  \\par Motivated by some of the uncertainty surrounding the most appropriate choice of initial conditions for problems of this type, we also carried out some idealised simulations in which the gas pressure was not adjusted when the flux tube was introduced. We stress that this is not a realistic situation for the solar convection zone given that this leads to a strong imbalance between the magnetic pressure and the local gas pressure. However, this does allow the system to adjust to the presence of the flux tube in a self-consistent way. This adjustment occurs very rapidly and (somewhat remarkably) the subsequent evolution of the system is comparable to that of the previous set of simulations, at least at comparable magnetic Reynolds numbers. Results from these idealised calculations strongly suggest that the only aspect of the initial conditions that influences the evolution of the flux tube is the initial magnetic field distribution. This is an encouraging result because it suggests that flux emergence simulations are largely insensitive to the precise choice of initial conditions, which is a matter of considerable uncertainty in idealised models of this type."
    },
    "1208/1208.2722_arXiv.txt": {
        "abstract": "The Australia Telescope Large Area Survey (ATLAS) has surveyed seven square degrees of sky around the Chandra Deep Field South (CDFS) and the European Large Area ISO Survey - South 1 (ELAIS-S1) fields at 1.4\\,GHz. ATLAS aims to reach a uniform sensitivity of $10\\,\\mu$Jy\\,beam$^{-1}$ rms over the entire region with data release 1 currently reaching $\\sim30\\,\\mu$Jy\\,beam$^{-1}$ rms. Here we present 466 new spectroscopic redshifts for radio sources in ATLAS as part of our optical follow-up program. Of the 466 radio sources with new spectroscopic redshifts, 142 have star-forming optical spectra, 282 show evidence for AGN in their optical spectra, 10 have stellar spectra and 32 have spectra revealing redshifts, but with insufficient features to classify. We compare our spectroscopic classifications with two mid-infrared diagnostics and find them to be in broad agreement. We also construct the radio luminosity function for star-forming galaxies to z~$= 0.5$ and for AGN to z~$= 0.8$. The radio luminosity function for star-forming galaxies appears to be in good agreement with previous studies. The radio luminosity function for AGN appears higher than previous studies of the local AGN radio luminosity function. We explore the possibility of evolution, cosmic variance and classification techniques affecting the AGN radio luminosity function. ATLAS is a pathfinder for the forthcoming EMU survey and the data presented in this paper will be used to guide EMU's survey design and early science papers. \\\\ ",
        "introduction": "\\label{intro} Extragalactic radio sources comprise both star-forming (SF) galaxies and active galactic nuclei (AGN) \\citep{Condon92}. In order to understand the history of the Universe we must unravel the cosmic evolution of both classes.  The galaxies that host radio sources display a range of different optical spectra. SF galaxy spectra are made up of direct emission from starlight and emission lines from H\\textsc{ii} regions, and are quite easily identified at low redshifts. Some radio-loud AGN galaxies display strong high-excitation lines, either broad or narrow depending on obscuration and the orientation \\citep{Antonucci93}, while others have weak or no emission features \\citep[e.g.][]{Sadler02}. These spectral differences have been attributed to different accretion modes \\citep{Hardcastle07, Croton06}. Sources without high-excitation lines are fuelled by the accretion of hot gas (`hot mode') while high-excitation sources require fuelling by cold gas (`cold-mode'). The radiatively inefficient `hot-mode' accretion is driven by the gravitational instability \\citep{Pope12}, and results in a geometrically thick, optically thin accretion disk within the hot broad-line region. On the other hand, the radiatively efficient `cold-mode' accretion results in a geometrically thin, optically thick accretion disk within the broad-line region, and a narrow-line region further out. `Cold-mode' accretion may be triggered by a merger event \\citep[e.g.][]{Shabala11}, which can also trigger star-formation within the galaxy. Consequently, sources undergoing `cold-mode' accretion may appear as `hybrid' objects, that is, objects whose radio emission results from both star-formation and an AGN component \\citep[e.g. PRONGS, ][]{Mao10b}. \\citet{Norris11b} find that the extreme ULIRG, F00183-7111, is such an object.  Early radio surveys such as the 3C and 3CR Radio Surveys \\citep{Edge59,Bennett62} had rms levels of a few Janskys and detected predominately powerful radio galaxies and quasars, allowing the cosmic history of these sources to be determined. Subsequent studies found that powerful radio galaxies and quasars evolve very strongly with redshift and the space density of these sources at z $\\sim$ 2 is orders of magnitude higher than in the local Universe \\citep[e.g.][]{Dunlop90}.  More recent wide-field radio surveys such as NVSS \\citep{Condon98}, SUMSS \\citep{Mauch03} and FIRST \\citep{Becker95} have reached rms levels of a few mJy. At these flux densities the radio source population is dominated by radio galaxies and quasars powered by AGN \\citep[e.g.][]{Sadler02}. The most powerful of these can be seen out to the edge of the visible Universe.  Source counts show a well-characterised upturn below 1\\,mJy\\,beam$^{-1}$, which is above that predicted from the extrapolation of source counts for AGN with high-flux densities. Some studies attribute this to strong evolution of the SF galaxy luminosity function \\citep[e.g.][]{Hopkins98} while others find a significant contribution from lower-luminosity AGN \\citep[e.g.][]{Huynh07}. \\citet{Sadler07} found that low-luminosity AGN undergo significant cosmic evolution out to z = 0.7, consistent with studies of the optical luminosity function for AGN \\citep[e.g.][]{Croom04}.  The deepest radio surveys now reach rms levels below 10\\,$\\mu$Jy\\,beam$^{-1}$ \\citep[e.g.][]{Owen08}, but cover very small areas on the sky. The distribution of AGN and SF galaxies at low flux densities is not well known, with some studies finding the proportion of AGN declining with decreasing flux density and emission due to star-formation (SF) dominating \\citep[e.g.][]{Mauch07, Seymour08}, while other studies find the distribution is closer to a 50/50 split between SF galaxies and AGN  down to $\\sim$50\\,$\\mu$Jy\\,beam$^{-1}$ \\citep[e.g.][]{Smolcic08, Padovani09}. The lack of consensus of the distribution of AGN and SF at low flux densities is due to the small areas of sky surveyed at these sensitivities. \\citet{Norris11a} provide a brief comparison of the various studies and, in their figure 4, show that the fraction of SF galaxies at $<$100\\,$\\mu$Jy\\,beam$^{-1}$ could range from 30 - 80 per cent.  To date, understanding the evolution of the faint radio source population has been limited by the availability of sufficiently deep wide-area radio surveys. With the aim of imaging approximately seven square degrees of sky over two fields to 10\\,$\\mu$Jy\\,beam$^{-1}$ at 1.4\\,GHz, the Australia Telescope Large Area Survey \\citep[ATLAS,][]{Norris06,Middelberg08} is the widest deep field radio survey yet attempted. ATLAS observes two separate fields so as to minimize cosmic variance \\citep{Moster11}: Chandra Deep Field South (CDFS) and European Large Area ISO Survey-South~1 (ELAIS-S1). The two fields were chosen to coincide with the \\textit{Spitzer} Wide-Area Infrared Extragalactic (SWIRE) Survey program \\citep{Lonsdale03}, so optical and mid-infrared identifications exist for most of the radio objects. CDFS also encompasses the Southern Great Observatories Origins Deep Survey (GOODS) field \\citep{Giavalisco04}.  The first data release (Data Release 1, or DR1) from ATLAS consists of the preliminary data published by \\citet{Norris06} and \\citet{Middelberg08} and reaches an rms sensitivity of $\\sim$30$\\,\\mu$Jy\\,beam$^{-1}$. The flux densities of DR1 are accurate to $\\sim$10 per cent. The final ATLAS data release (DR3: Banfield et al., in preparation) will include additional data taken with the new CABB upgrade \\citep{Wilson11} to the ATCA, and is expected to reach an rms sensitivity of $\\sim$10$\\,\\mu$Jy\\,beam$^{-1}$. For this paper we use radio data only from DR1.  Spectroscopic data provide accurate redshift information that can be used to determine absolute magnitudes, luminosities etc. This paper presents spectroscopic redshifts and spectroscopic classifications of radio sources in the ATLAS fields obtained using the AAOmega multi-fibre spectrograph at the Anglo-Australian Telescope (AAT) of sources in the ATLAS fields. Section 2 describes the observations and data analysis while Section 3 provides spectroscopic classifications and a description of the redshift catalogue. Section 4 discusses how the spectroscopic classifications compare to mid-infrared diagnostics and presents the radio luminosity function for ELAIS. Section 5 summarises our conclusions.  This paper uses H$_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega$$_M$ = 0.3 and $\\Omega$$_\\Lambda$ = 0.7 and the web-based calculator of \\citet{Wright06} to estimate the physical parameters. Vega magnitudes are used throughout. ",
        "conclusions": "We have presented 466 new spectroscopic redshifts for radio sources in ELAIS and CDFS as part of ATLAS, using AAOmega on the AAT. We have classified the sources as AGN or SF based on their spectra and compared these spectroscopic classifications with mid-IR diagnostics and found them to be in good agreement. Of the 466 sources, 282 are AGN, 142 are SF, 32 are either SF or AGN (none of the diagnostics used in this paper were able to determine if they were SF or AGN) and 10 are chance alignments with stars.  We have constructed the RLF for ELAIS for both SF galaxies and AGN. We find positive evolution, consistent with previous studies, to z~=~0.5 for the SF RLF. We find our AGN RLF to be consistent with previous work by \\citet{Padovani11}, and a factor of $\\sim$3 higher than the work by \\citet{Mauch07}.  We cannot make any definitive statements as to why the SF RLF for ELAIS is consistent with previous studies whereas the AGN RLF for ELAIS appears inconsistent. We attribute the inconsistencies to the possibility of evolution, cosmic variance, obtaining spectra of only the bulge of a SF galaxy leading to an overestimate of early-type galaxies, or possibly an amalgamation of all these factors. We shall perform these analyses on DR3 where the extra data will allow us to draw more definitive conclusions.   ATLAS is the pathfinder for the forthcoming EMU Survey \\citep{Norris11a}, planned for the new ASKAP telescope \\citep{Johnston08}, which will survey the entire visible sky to an rms depth of 10$\\,\\mu$Jy\\,beam$^{-1}$. Because the ATLAS and EMU surveys are well-matched in sensitivity and resolution, the results obtained in this and other papers on the ATLAS survey will be used to guide the survey design and early science papers for EMU. In particular, the spectroscopy on ATLAS sources will provide a valuable training set to guide the algorithms used to determine the redshift distribution of the anticipated 70 million EMU sources."
    },
    "1208/1208.0050.txt": {
        "abstract": "The Atacama Cosmology Telescope (ACT) is a six-meter microwave telescope for measuring arc minute scale anisotropy of the Cosmic Microwave Background (CMB). After its first light in October 2007, two seasons of observations have been carried out with the Millimeter Bolometric Array Camera (MBAC), currently composed by three, 32 by 32 element, Transition Edge Sensor (TES) bolometer arrays, for observations at 145~GHz, 220~GHz and 280~GHz respectively. For generating the CMB maps, a suite of sophisticated analysis techniques have been developed and implemented to produce a customized data pipeline. This paper provides the full characterization of the data as obtained from the telescope, including noise and systematic signals, together with the set of preprocessing steps and tools currently used for map-making. ",
        "introduction": "\\label{sec_introduction} Brief description of: \\begin{slim_itemize} \\item{The science goals} \\item{The site} \\item{The telescope, camera (MBAC), detectors and readout circuit. Reference to D.~Swetz~et~al.} \\item{Seasons observations including number of observed hours in both stripes.} \\item{Description of the data set, stressing its size.} \\item{Outline of the paper} \\end{slim_itemize} ",
        "conclusions": "\\label{sec_conclusions}  \\begin{slim_itemize} \\item{Step by step summary of the data pipeline from raw to maps.} \\item{Summary of noise properties, including brief comments about the integration time required for our science goals.} \\item{Happy ending.} \\end{slim_itemize}  {\\em Acknowledgements.}   %"
    },
    "1208/1208.5102_arXiv.txt": {
        "abstract": "Dynamical Chern-Simons gravity cannot be strongly constrained with current experiments because it reduces to General Relativity in the weak-field limit. This theory, however, introduces modifications in the non-linear, dynamical regime, and thus, it could be greatly constrained with gravitational waves from the late inspiral of black hole binaries. We complete the first self-consistent calculation of such gravitational waves in this theory. For favorable spin-orientations, advanced ground-based detectors may improve existing solar-system constraints by 6 orders of magnitude. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.5272_arXiv.txt": {
        "abstract": "We propose a plausible mechanism to explain the formation of the so-called ``obscuring tori'' around active galactic nuclei (AGNs) based on three-dimensional hydrodynamic simulations including radiative feedback from the central source. The X-ray heating and radiation pressure on the gas are explicitly calculated using a ray-tracing method. This radiation feedback drives a ``fountain'', that is, a vertical circulation of gas in the central a few to tens parsecs. Interaction between the non-steady outflows and inflows causes the formation of a geometrically thick torus with internal turbulent motion. As a result, the AGN is obscured for a wide range of solid angles. In a quasi-steady state, the opening angles for the column density toward a black hole $< 10^{23}$ cm$^{-2}$ are approximately $\\pm 30^{\\circ}$ and $\\pm 50^{\\circ}$ for AGNs with 10\\% and 1\\% Eddington luminosity, respectively. Mass inflows through the torus coexist with the outflow and internal turbulent motion, and the average mass accretion rate to the central parsec region is $2\\times 10^{-4}\\sim 10^{-3}\\, M_{\\odot}\\; {\\rm yr}^{-1}$; this is about ten times smaller than accretion rate required to maintain the AGN luminosity. This implies that relatively luminous AGN activity is intrinsically intermittent or that there are other mechanisms, such as stellar energy feedback, that enhance the mass accretion to the center. ",
        "introduction": "The presence of optically thick obscuring tori has been postulated to explain the various observed properties of active galactic nuclei (AGNs), particularly with respect to the two major categories of AGNs, namely, type 1 and type 2 \\citep{antonucci93}.  However, the origin of the optically and geometrically thick material that covers a large solid angle with respect to the AGN and the nature of its three-dimensional structure are still unclear, although its geometry is often schematically described as a ``doughnut-shaped'' \\citep[e.g.,][]{urry95}. In fact, statistical studies of the absorption features of local AGNs suggested have suggested the presence of complex, inhomogeneous obscuring material around the central engine \\citep{shi2006, hicks09}. The clumpy torus model is also consistent with the observed spectral energy distributions (SEDs) \\citep[e.g.,][]{nenkova2002,elitzur06,levenson2009,schartmann09, stalevski2012}. The essential question that then arises is regarding how such an inhomogeneous `torus' is formed and how it attains such a large covering fraction against the energy dissipation of the clumpy medium in the gravitational field of the central black hole.  Numerous theoretical models have been proposed to explain the origin of the obscuring material. On an accretion disk scale the proposed causes include hydromagnetic disk winds \\citep{konigl1994}, disk-driven hydromagnetic winds and radiation pressure due to continuum emission from the accretion disk \\citep{keating2012}, and warping of irradiated accretion disks \\citep{arm97}. Although a pasec-scale dusty torus is observed in some nearby AGNs \\citep{jaffe04}, the nucleus could be obscured by relatively less dense ISM in the range of  10-100 pc or even by galactic disks themselves \\citep{levenson2001,goulding2012}. There are many instances where evidence suggests that most AGNs are associated with circumnuclear starbursts \\citep[e.g.][]{imanishi04, davies07, hicks09,chen09, imanishi2011, Diamond2012, Woo2012}. It is suggested that the star formation activities drive mass inflow into AGNs, thereby leading to the growth of black holes  \\citep[see the review by][]{alexander2011}. The nuclear starbursts themselves could inflate the circumnuclear disk and obscure the central source. \\citet{wada02} and \\citet{wada09} have proposed that a clumpy torus-like structure is naturally reproduced due to energy feedback from supernova explosions (SNe) in a disk. In their model, the internal density and temperature of the thick disk are highly inhomogeneous, and the velocity field of the disk is turbulent. The scale height of the disk is determined by the balance between the turbulent energy dissipation and energy input from SNe. More recently, \\citet{hopkins2012} have proposed that 10-pc-scale thick tori can be formed even without the existence of strong stellar energy feedback if gas inflow from kiloparsec scales drives instability in the circumnuclear region.  On a parsec scale or smaller, geometrically and optically thick torus sustained by radiation pressure have been suggested \\citep{kro88,pie92}. Following this model, static solutions of thick disks supported by infrared radiation have been explored by \\citet{krolik07} and \\citet{shi-krolik2008}. \\citet{ohs01} have suggested that static, geometrically thin obscuring walls could be formed by the interplay between the radiation from the nucleus and a ring-like nuclear starburst.   It would be more natural to assume that radiative feedback causes the dynamical evolution of the surrounding ISM rather than static structures. \\citet{roth2012} calculated the net rate of momentum deposition due to the radiation pressure in the surrounding gas and estimated the mass-loss rate by outflow using Monte Carlo radiative transfer calculations. \\citet{dorodnitsyn2012} performed `2.5D' (i.e. basic equations are solved in three dimensions with axisymmetry) radiative hydrodynamic simulations, and they found that an AGN torus can be better described in terms of a radiationally supported flow rather than a quasi-static obscuring torus. Using 2D simulations, \\citet{schartmann2011} studied the evolution of dusty clouds irradiated by the AGN and found that the radiative transfer effect is significant depending on the column density of the clouds.  As is usual in multi-dimensional radiation-hydrodynamic simulations, cirtain simplifications have been applied in these previous studies, such as assuming flux-limited diffusion approximation and introducing symmetry for the rotational axis and equatorial plane of the AGN. In this study, we examine the effect of the radiation from the central source on the ISM by assuming a different set of simplifications. We performed fully three-dimensional hydrodynamic simulations on scale of a few tens pc around an AGN without introducing any symmetries. We took into account radiative heating and pressure due to ``direct radiation'' from the central source along rays for all the $256^{3}$ grid cells; however, we ignored the phenomenon of scattering/re-emission. This is a natural extension of our previous three-dimensional hydrodynamic simulations of tens-of-parsec scale tori \\citep{wada02,wada05, yamada07,wada09}, in which self-gravity of the ISM, radiative cooling, SNe, and UV heating were taken into account. Here, we report the first results without SNe feedback. The SED analysis based on the present model will appear elsewhere (Schartmann and Wada, in prep.). In this study, we propose a dynamical process leading to the formation of thick tori, and we examine their structures in a quasi-steady state. Further, we discuss the changes in the column density and the net mass accretion rate depending on the luminosity of the AGN. ",
        "conclusions": "In our study, we showed that a geometrically and optically thick torus can be naturally formed in the central region extending tens parsecs around a low luminosity AGN. We found that radiation drives a ``fountain'' of gas, i.e. vertical circulation, and this also naturally produces a thick, turbulent, and inhomogeneous disk that resembles a torus. The opening angle of the torus for  $N < 10^{23}$ cm$^{-3}$ is $60-100^{\\circ}$ for $L_{AGN} = 0.1-0.01 L_{E}$, in contrast to 140$^{\\circ}$ for the model without radiative feedback. The average accretion rate to the central parsec region is larger in a less luminous AGN; however, this accretion rate is insufficient to maintain the AGN luminosity over several Myrs. These results imply that the AGN luminosity and structure of the surrounding ISM alternate between an active phase (high luminosity and a torus with a small opening angle) and inactive phase (low luminosity and a torus with a large opening angle). For thick tori (Fig. \\ref{wada_fig: f4}) to maintained their structures for several million years or longer, there should be other mechanisms to enhance the accretion rate, such as supernova-driven turbulence \\citep{wada02,wada09}, stellar mass loss \\citep{schartmann2010}, turbulence drives by galactic inflows \\citep{hopkins2012} and radiation drag \\citep{kawakatu2002}. In fact, recent observations suggest a tight correlation between star formation rate and the Eddington ratio of the central BH \\citep{chen09}.  Figure \\ref{wada_fig: f5} shows that the covering fraction of the central source is smaller in AGNs with lower luminosity. This is because more mass can be ejected from the central region to higher latitudes, and as a result, the radiation from the AGN is self-obscured. Therefore, the radiation-driven outflows appear only in a small solid angle around the rotational axis, as explained in \\S 3.1 and illustrated in Fig. \\ref{wada_fig: f2b}. This means that more luminous AGNs are obscured over larger solid angles, although this may appear inconsistent with the observed ``receding torus'' \\citep{lawrence1991, ueda2003}. For example, \\citet{hasinger2008} studied 1290 AGNs and found that the absorbed fraction {\\it decreases} with X-ray luminosity in the range of $L_{X}=10^{42-46}$ erg s$^{-1}$. However, it is to be noted that the AGNs explored in the present simulation have relatively low luminosity, i.e., in the range of $L_{X}=10^{42-43}$ erg s$^{-1}$. In this small luminosity range, the structure of the torus does not vary significantly, especially for higher values of column density ($N > 10^{24}$ cm$^{-2}$ ). For a central source that is 100 times brighter, the radiation-dominant region is considerably larger, as shown by the dashed line in Fig. \\ref{wada_fig: f2b}, and therefore, we can expect a larger opening angle for this case. More recently, \\citet{lu2010} have reportred results that are significantly different from those of previous studies; in their SDSS/FIRST survey,  they found that the type-1 fraction exhibits a constant of ~20\\% in the [O III] 5007 luminosity range of $40.7< \\log(L_{[{\\rm O}_{\\rm III}]}/ {\\rm erg}\\, {\\rm s}^{-1}) <43.5$. They also found that the type-1 fraction is independent of the Eddington ratio for its value between 0.01 and 1, if only high density ($> 10 M_{\\odot}\\; {\\rm pc}^{-3}$) gas in the torus contributes to the obscuration; this is also the case as seen in Fig. \\ref{wada_fig: f4}.  {Further, this study shows that the global structures of the torus and outflows are not static, and there is never a perfect symmetry in terms of the equatorial plane and the rotation axis  (Fig. \\ref{wada_fig: rho-azimuth}), even if we assume the presence of angle-dependent, axisymmetric flux from the central source. The opening angle of the torus and other structures could be affected by the choice of the central radiation source. In this study, we assumed a thin accretion disk; however, recent two-dimensional radiation-magnetohydrodynamic simulations have shown that geometrically thick and radiatively inefficient accretion flows associated with outflows are formed depending on the central gas density \\citep{ohs11}. In this case, the radiation flux is collimated around the rotation axis, and the effect of the radiation on the disk might be less than that observed in the models used in this study. Even for a thin accretion disk, the rotational axis is not necessarily aligned to the galactic rotational axis. This causes anisotropic illumination of the ISM, and this misalignment could also affect the occultation of the central source by the torus \\citep[see also][]{kawaguchi2011}. Consequently, we can expect the formation of more such anisotropic structures of the torus and outflows in such cases. This can be confirmed by high-resolution observations of nuclear active galaxies associated with molecular outflows, such as NGC 1377 \\citep{aalto2012b} and by ALMA in the near future.}    \\vspace{0.5cm} The author is grateful to M. Schartmann for providing the dust absorption tables and also for his helpful comments regarding the draft. We also thank M. Spaans and the anonymous referee for many valuable comments. The numerical computations presented in this paper were carried out on the NEC SX-9 in the Center for Computational Astrophysics, NAOJ and Cyberscience Center, Tohoku University. The author thanks N. Horiuchi, Y. Sakaguchi, and the NEC Corporation for optimizing the code for the SX-9. This work is partly supported by Grant-in-Aid for Scientific Research (C) 23540267."
    },
    "1208/1208.0958_arXiv.txt": {
        "abstract": "While the importance of dusty asymptotic giant branch (AGB) stars to galactic chemical enrichment is widely recognised, a sophisticated understanding of the dust formation and wind-driving mechanisms has proven elusive due in part to the difficulty in spatially-resolving the dust formation regions themselves. We have observed twenty dust-enshrouded AGB stars as part of the Keck Aperture Masking Experiment, resolving all of them in multiple near-infrared bands between 1.5$\\mu$m and 3.1$\\mu$m. We find 45\\% of the targets to show measurable elongations that, when correcting for the greater distances of the targets, would correspond to significantly asymmetric dust shells on par with the well-known cases of IRC~+10216 or CIT~6. Using radiative transfer models, we find the sublimation temperature of $T_{\\rm sub}$(silicates) $=1130\\pm90$\\,K and $T_{\\rm sub}$(amorphous carbon) $=1170\\pm60$\\,K, both somewhat lower than expected from laboratory measurements and vastly below temperatures inferred from the inner edge of YSO disks.   The fact that O-rich and C-rich dust types showed the same sublimation temperature was surprising as well. For the most optically-thick shells ($\\tau_{2.2\\mu{\\rm m}}>2$), the  temperature profile of the inner dust shell is observed to change substantially, an effect we suggest could arise when individual dust clumps  become optically-thick at the highest mass-loss rates. ",
        "introduction": "One of the most dramatic phases in the life of an intermediate mass star is the Asymptotic Giant Branch, a relatively short period where a star loses most of its initial mass through a dusty wind.  Researchers still do not understand all the ingredients necessary for producing the high mass-loss rates observed during this stage.  The massive envelopes ejected during this phase are thought to be later illuminated during the planetary nebula stage, a stage where most stars show strong bipolar circumstellar structures \\citep{balick2002}.  Following the advent of infrared detectors, early workers made simple spherically-symmetric models of dusty shells around large samples of AGB stars fitting only to the spectral energy distributions\\citep[e.g.,][]{rrh1982,rrh1983,rrh1983b}.  M-type stars are typically surrounded by dust shells composed of amorphous silicates while C-stars have carbonaceous dust. These early workers were able to show that dust condensed around 1000~K within a few stellar radii of the stars and also estimated mass-loss rates typically 10$^{-6}$\\msun/yr and as high as 10$^{-4}$\\msun/yr.  More recently, \\citet{ivezic1995} developed the code DUSTY to study dust shells in a systematic way and made models for a large sample of stars, again fitting just the spectral energy distribution.  The simple picture of spherically-symmetric and uniform mass-loss was challenged by the observations of the Infrared Spatial Interferometer (ISI), a long-baseline mid-infrared interferometer  \\citep{danchi1994}. These workers found a diversity of shell morphologies with some red giants showing episodic dust shells ejections and others with a more continuous distribution of dust. A more dynamic and asymmetric vision of mass-loss fit into debates into the origins of bipolar symmetry in planetary nebulae.  High angular resolution near-infrared speckle and aperture masking on 8-m class telescopes were able to image fine details on some dust shells, such as the prototype carbon star IRC~+10216 \\citep[e.g.][]{tuthill2000}.  An elaborate model was presented by \\citet{menshchikov2002} arguing for complex, spatially-varying dust properties and density structures.  While IRC~+10216 shows complexity within the inner few stellar radii, it is unclear if these structures represent {\\em global} asymmetries or just {\\em weather conditions} of the dust formation process observed {\\em in situ}.  Here we present the full dataset of dust-enshrouded giants observed with the 10-year project called the Keck Aperture Masking Experiment \\citep{tuthill2000b}.  This experiment delivered well-calibrated spatial information on the scale of $\\sim$50~milliarcseconds (mas) in the astronomical K band ($\\lambda_0=2.2\\mu$m), enough to resolve all the dusty targets presented here and to measure their dust shell sizes and asymmetries.  This paper includes 20 objects with observations in typically 3 wavelengths ranges, 1.65$\\mu$m, 2.2$\\mu$m, and 3.1$\\mu$m. We have also extracted photometry to construct coeval near-IR spectral energy distributions -- an important factor since these objects pulsate and show large variations in flux on yearly timescales. Lastly, we used a radiative transfer code to fit each epoch of each target star using simultaneously the NIR photometry and multi-wavelength angular size information from Keck masking.  The primary goals of these observations and modeling efforts are to measure the physical characteristics of a large sample of the most extreme dusty AGB stars, to address the question of the onset of circumstellar asymmetries, to determine any differences between silicate and carbon-rich dust shells, and to constrain the optical properties of the dust particles themselves.  Lastly, this publication marks the final large data release of AGB star data from our diffraction-limited Keck masking experiment and we anticipate this work will provide a rich dataset for more detailed modelling efforts by other workers. ",
        "conclusions": "\\label{discussion} Our survey provides the first constraints on the asymmetry of the dust shells for such a large sample of  dust-enshrouded AGB stars. We found that 4 out of 7 M-stars and 5 of 13 C-stars showed evidence of dust shell asymmetries, with dust shell elongations between 10\\% and 40\\%. While this level of asymmetry may sound mild, it actually (quantitatively) compares to the level of asymmetry that would be expected for the most asymmetric dust shells known if placed at 1~kpc.  For instance, we know that IRC~+10216 \\citep{tuthill2000} and CIT~6 \\citep{monnier2000} have dramatic global asymmetries in their dust shell, detailed imaging made possible by virtue of their proximity. If we placed these targets farther away, we would not be able to image the detail but they would appear $\\sim$20\\% elongated, similar to the degree observed here in 45\\% of our sample.  For CIT~3, we confirm the asymmetries seen by \\citet{hofmann2001} and note that  \\citet{vinkovic2004} showed that the 20\\% elongation could be explained by a bipolar outflow. That said, clumpy dust formation \\citep{fleischer1992} might also cause stochastic variations in the inner dust shell geometry that could appear as short-lived elongations.   Mid-infrared observations with long-baseline interferometers (e.g, ISI, VLTI-MIDI) should focus on these targets to determine the nature of the asymmetries.  In addition, long-term monitoring of these dust shells will help settle debates concerning when the environments of evolved stars develop large scale asymmetries commonly revealed in the later planetary nebula stage.  For instance, a long-term asymmetry in a constant position angle (as judged by linear polarization or spatially resolved data) would be a sign of a global bipolar mass-loss asymmetry and not just {\\em weather}.  In order to look at dust shell properties for our full sample, we have plotted the inner edge dust temperature $T_{\\rm dust}$ vs total dust shell optical depth $\\tau_{2.2\\mu\\,m}$ for all our targets.  Figure 21 shows these results split into O-rich and C-rich dust type.   For K-band optical depths below 2, we find the sublimation temperature of $T_{\\rm sub}$(silicates) $=1130\\pm90$\\,K and $T_{\\rm sub}$(amorphous carbon) $=1170\\pm60$\\,K, both somewhat lower than expected from laboratory measurements \\citep{lodders1999} and vastly below temperatures inferred from the inner edge of YSO disks \\citep[$\\sim$1800K,][]{tannirkulam2008, benisty2010}. One component to the observed lower dust temperature could be due to the fact that the central star varies in luminosity by about a factor of 2 during the pulsation cycle and we see the dust cooler than the condensation temperature during phases away from maximum light.  The $T_{\\rm dust}$ vs optical depth $\\tau_{2.2\\mu\\,m}$ diagram (Figure~21) also shows  no statistically-significant difference between O-rich and C-rich dust types,  counter to expectation of higher temperatures for carbon-rich dust  \\citep{lodders1999}.   We recognise that our simple dust shell modeling may not lead to accurate estimates of the dust sublimation temperature if the inner dust formation environment radically departs from a power law density distribution, perhaps due to pulsations, timescale for dust formation, or multiple dust species.  Interestingly though these concerns would likely affect C-rich and O-rich shells similarly and so the lack of a clear difference in sublimation temperatures between these dust types appears robust.    The other important feature of Figure 21, $T_{\\rm dust}$ vs optical depth $\\tau_{2.2\\mu\\,m}$, is the apparent temperature at the inner edge of the dust shell gets lower and lower with increasing optical depths above 2.  This appears true for both C-rich and O-rich shells.  Here we do not believe we are seeing an actual reduction in the dust sublimation temperature, but rather a change in the temperature profile in the inner dust formation zone due to a breakdown in the assumption of a spherically-symmetric $r^{-2}$ density power law. We have ample evidence that dust formation is clumpy, as has been imaged in great detail for IRC +10216 \\citep{tuthill2000}, but these clumps have been shown to have a relatively weak affect on the temperature structure for low optical depths.  Next we further explore how a clumpy dusty environment could change the temperature profile of the dust shell when the individual clumps become themselves optically thick to the stellar and even hot dust radiation field.  Clumpy structures are seen to evolve in 2-D models of dust shells due to self-amplifying density perturbations \\citep[e.g.][]{Woitke2000}.   First optically thick dust regions form and these regions cast shadows on the dust behind them.  Consequently, the temperatures decrease by 100's of degrees K and this allows for a higher rate of dust formation in these shadow regions.  Scattering and re-emission of light by the optical thick regions increases the intensity of radiation between them and eventually the light escapes through the optically thin regions in between the optically thick regions.  Thereupon, the temperature within the optically thin regions increases, which decreases the rate of dust production.  These processes thus amplify the initial homogeneities until large-scale clumpy structures start to form, such as ``dust fingers\" \\citep{Woitke2005}.  Indeed, \\citet{Woitke2005} did see {\\em average dust temperatures to be reduced} due to these opacity effects but at much weaker level than we see in Figure~21.  Realizing that our data reveal a strong effect only at $\\tau$'s several times larger than probed by \\citet{Woitke2005}, we suggest that dust shadowing effects get dramatically stronger  when individual clumps become optically thick to both stellar radiation as well as hot dust emission. A 3D radiative transfer calculation of a dusty dust shell could validate or disprove this explanation.  In conclusion, our large sample of  spatially resolved dust-enshrouded stars have led to new insights into the late stages of AGB star evolution. We find levels of dust shell elongations that point to significant asymmetries in nearly half of our targets.  Our spatial and SED data combined has eliminated some model degeneracies, and we now have the best constraints on the actual sublimation temperatures for dust forming in this outflows, finding lower temperatures than expected from terrestrial experiments and not confirming the large difference expected between carbon-rich and silicate-rich dust. Lastly, we discovered a systematic change in the temperature profile for inner-most dust regions when the dust shell optical depth rises above $\\tau_{2.2\\mu\\,m}>2$.  This observed lowering of the central dust temperatures could be naturally explained as a consequence of shadowing caused by clumpy dust formation on spatial scales smaller than our angular resolution, but other possibilities should be further explored as well."
    },
    "1208/1208.5050_arXiv.txt": {
        "abstract": "We study the comoving space density of X-ray-selected luminous active galactic nuclei (AGNs) and the obscured AGN fraction at high redshifts ($3 < z < 5$) in the Subaru/{\\it XMM-Newton} Deep Survey (SXDS) field. From an X-ray source catalog with high completeness of optical identification thanks to deep optical images, we select a sample of 30 AGNs at $z > 3$ with intrinsic (de-absorbed and rest-frame 2--10 keV) luminosities of $L_{\\rm X} = 10^{44-45}$ erg s$^{-1}$ detected in the 0.5--2 keV band, consisting of 20 and 10 objects with spectroscopic and photometric redshifts, respectively. Utilizing the $1/V_{\\rm max}$ method, we confirm that the comoving space density of luminous AGNs decreases with redshift above $z > 3$. When combined with the {\\it Chandra}-COSMOS result of Civano et al.\\ (2011), the density decline of AGNs with $L_{\\rm X} = 10^{44-45}$ erg s$^{-1}$ is well represented by a power law of $(1 + z)^{-6.2 \\pm 0.9}$. We also determine the fraction of X-ray obscured AGNs with $N_{\\rm H} > 10^{22}$ cm$^{-2}$ in the Compton-thin population to be 0.54$^{+0.17}_{-0.19}$, by carefully taking into account observational biases including the effects of photon statistics for each source. This result is consistent with an independent determination of the type-2 AGN fraction based on optical properties, for which the fraction is found to be 0.59$\\pm$0.09. Comparing our result with that obtained in the local Universe, we conclude that the obscured fraction of luminous AGNs increases significantly from $z=0$ to $z>3$ by a factor of 2.5$\\pm$1.1. ",
        "introduction": "The evolution of active galactic nuclei (AGNs) carries key information for understanding the growth history of supermassive black holes (SMBHs) in galactic centers. Recent X-ray surveys have revealed that luminous AGNs have their number density peak at higher redshifts, $z \\sim 2-3$, than less luminous ones \\citep{ueda2003, barger2005, lafranca2005, hasinger2005, silverman2008, ebrero2009, yencho2009, aird2010}. This behavior is called ``down-sizing'' or ``anti-hierarchical'' evolution.  These results suggest that present-day massive SMBHs formed in the early epoch of the Universe, compared with less massive ones. A similar evolution is also observed in the star forming history of galaxies \\citep[e.g.,][]{cowie1996, kodama2004}. This provides a further evidence for the cosmological ``co-evolution'' scenario of SMBHs and galaxies, which is expected from the tight correlation between a SMBH mass and galactic bulge properties observed in the local Universe \\citep[see e.g.,][]{magorrian1998, ferrarese2000, gebhardt2000}.  To elucidate the formation processes of SMBHs, it is very important to reveal the AGN evolution at the high redshifts before the density peak, where the rapid growth of SMBHs took place. Optical/infrared AGN surveys, such as the Sloan Digital Sky Survey, the Canada-France High-z Quasar Survey, and the UKIDSS Large Area Survey, are able to detect very luminous quasars at high redshifts, now up to $z \\simeq 7$ \\citep{jiang2009, willott2010, mortlock2011}, revealing that their space number density has a peak around $z \\simeq 2-3$ and shows a rapid decline by 96\\% from $z = 3$ to $z = 6$ \\citep{richards2006}. Recently, the space density of lower luminosity AGNs has been also investigated from deep optical survey fields \\citep{glikman2011, ikeda2011}, although the results are still controversial due to the complexity of the corrections required for incompleteness and contamination. While these optical surveys are useful in detecting unobscured (so-called type-1) AGNs, they inevitably miss the population of obscured (type-2) AGNs, the major class of the AGN population \\citep[see e.g.,][]{gilli2007}.  Hard X-ray observations at rest-frame energies above a few keV provide an ideal opportunity to detect obscured AGNs, thanks to their much smaller biases against absorption than the optical band and low contamination from stars in the host galaxies. \\cite{silverman2008} reports that the space density of X-ray-selected luminous AGNs significantly declines at $z > 3$, similar to that of optically-selected QSOs, although they have to assume a completeness correction because the spectroscopic identification rate of the sample is not high ($\\sim50\\%$). More recently, \\cite{brusa2009} and \\cite{civano2011} have studied X-ray detected AGNs at $z \\simgt 3$ from the COSMOS survey \\citep{scoville2007, hasinger2007, elvis2009}, one of wide and deep multi-wavelength surveys over a continuous area. A decline of the space density above $z \\sim 3$ is confirmed for luminous AGNs with intrinsic 2--10 keV luminosities of $\\simgt 10^{44}$ erg s$^{-1}$. While the functional shape of this decline is found to be consistent with that of the exponential decline parameterization proposed by \\cite{schmidt1995}, which is based on the optically-selected QSOs, it still has a large uncertainty due to the limited sample size. For quantitative comparison with the results obtained in other wavelengths and theoretical models, it is important to establish the shape of the decline from a larger X-ray-selected AGN sample with high completeness.  Another key observational property of AGNs is the type-2 AGN fraction (or more precisely, the distribution of absorption column density), which gives strong constraints on the environments around SMBHs including the obscuring tori and host galaxies. According to the unified model of AGNs \\citep{antonucci1993, urry1995}, the classification of AGNs depends only on the viewing angle between the observer and the rotation axis of the torus. It is well known that the type-2 AGN fraction decreases with increasing X-ray luminosity \\citep{ueda2003, steffen2003, lafranca2005, hasinger2008, ueda2011}, indicating that the simplest unified model with a fixed geometry of tori does not hold. Several authors \\citep[e.g.,][]{lafranca2005, treister2006, hasinger2008} have furthermore reported that the type-2 fraction increases with redshift, suggesting an evolution of the structures around SMBHs. However, due to difficulties of type classification using photon-limited X-ray data and complex biases against obscuration, the intrinsic type-2 AGN fraction at $z>3$ has been only poorly investigated.  In this paper, we investigate the evolution of the comoving space density of X-ray-selected luminous AGNs and type-2 fraction at high redshifts ($3 < z < 5$) using the X-ray catalog of Subaru/{\\it XMM-Newton} Deep Survey (SXDS), another wide and deep multi-wavelength survey. The organization of the paper is as follows. Section~\\ref{sec:SAMPLE} describes the high-redshift AGN sample used in this study. In section~\\ref{sec:ANALYSIS}, we explain the AGN type classification methods using X-ray or optical data. The results of the comoving space density and type-2 fraction are presented in section~\\ref{sec:RESULT AND DISCUSSION}. Section~\\ref{sec:SUMMARY} summarizes the conclusions. Throughout this paper, the cosmological parameters ($H_{\\rm 0}$, $\\Omega_{\\rm M}$, $\\Omega_{\\Lambda}$) = (70 km s$^{-1}$ Mpc$^{-1}$, 0.3, 0.7) are adopted. The ``log'' symbol represents the base-10 logarithm. Quoted errors denote those at $1\\sigma$ level. ",
        "conclusions": "\\label{sec:SUMMARY}  Thanks to the deep and wide area coverage of the SXDS survey with extensive multi-wavelength follow-up observations, we are able to construct a highly complete X-ray-selected AGN sample at high redshifts ($3 < z < 5$), consisting of 20 and 10 AGNs with spectroscopic and photometric redshifts, respectively. The conclusions are summarized as follows:  \\begin{itemize}  \\item We derive the comoving space density of Compton-thin AGNs at $z$ = 3--5 with intrinsic luminosities of $L_{\\rm X}$ = 10$^{44-45}$ erg s$^{-1}$ at $z > 3$, using the $1/V_{\\rm max}$ method. We confirm the trend that the space density significantly declines with redshift as reported in previous works. Combining our SXDS result with that from the COSMOS survey \\citep{civano2011}, we find that the comoving space density decreases as $(1 + z)^{-6.2\\pm0.9}$ or $10^{-(0.60\\pm0.08)\\times(z-2.7)}$, providing so-far the best constraint on the declining profile of X-ray-selected AGNs at $z > 3$.  \\item By carefully taking into account observational biases, including statistical uncertainties in the absorption column densities derived from the X-ray data, we derive the obscured AGN fraction with $N_{\\rm H}> 10^{22}$ cm$^{-2}$ among those with $N_{\\rm H} <10^{24}$ cm$^{-2}$ to be 0.54$^{+0.17}_{-0.19}$ at $z$ = 3--5 in the luminosity range of $L_{\\rm X}$ = 10$^{44-45}$ erg s$^{-1}$. From comparison with the results obtained at lower redshifts, we establish that there is a significant redshift evolution of obscured fraction from $z = 0$ to $z > 3$ by a factor of $2.5\\pm1.1$ in the high-luminosity range.  \\end{itemize}"
    },
    "1208/1208.0169_arXiv.txt": {
        "abstract": "{}{}{}{}{} \\abstract {Resonant absorption  cyclotron features are a key  diagnostic tool to directly  measure the strength  of the  magnetic field  of accreting neutron stars.   However, typical values for  cyclotron features lie in the high-energy  part of the spectrum between 20  keV and 50 keV, where detection  is often damped  by the low statistics  from single pointed observations.} {We show that long-term monitoring campaign performed with Swift-BAT of persistently, but faint, accreting high-mass X-ray binaries  is able to reveal in their spectra the  presence of  cyclotron features.} {We extracted the average Swift-BAT 15-150 keV spectrum from the 54 months long Swift-BAT survey of the high-mass X-ray source IGR~J16493$-$4348. To constrain the broadband spectrum we used soft X-ray spectra from Swift-XRT and Suzaku pointed observations.} { We model the spectra using a set of phenomenological models usually adopted to describe the energy spectrum of accreting high-mass X-ray binaries; irrespective of the models we used, we found significant improvements in the spectral fits adding to the models a broad (10 keV width) absorption feature, with best-fitting energy estimate between 30 and 33 keV, that we interpret as evidence for a  resonant cyclotron absorption feature. We also discuss instrumental issues related to the use of Swift-BAT for this kind of studies and the statistical method to weight the confidence level of this detection. Correcting  for  the  gravitational  redshift of  a  1.4  M$_{\\sun}$ neutron star,  the inferred surface magnetic field is $B_{\\rm surf} \\sim 3.7  \\times 10^{12}$  Gauss. The spectral parameters of IGR~J16493$-$4348 fit well with empirical correlations observed when the whole sample of high-mass binaries with detected cyclotron features is considered.} {} ",
        "introduction": "} The Burst Alert Telescope  (BAT, \\citealp{barthelmy05}) on board Swift \\citep{gehrels04} has  been performing a continuous  monitoring of the sky in the hard X-ray  energy range (15--150 keV) since November 2004. The telescope, thanks  to its large field of  view (1.4 steradian half coded) and its  pointing strategy, covers a fraction  between 50\\% and 80\\% of the sky every day.   This has allowed the detection of many of the    new   INTEGRAL    High-Mass   X-ray    Binaries   \\citep[HMXBs, e.g.][]{cusumano10b}  and the  collection of  their long  term light curves  and spectra.   The  long and  continuous  monitoring of  these sources allows  to investigate the intrinsic  emission variability, to search for  long periodicities (orbital  periods) and to  discover the presence of  eclipse events.  The role of  Swift-BAT (hereafter simply BAT) is therefore fundamental to unveil the nature and the geometry of these binary  systems. Moreover BAT collects  their long-term averaged energy spectra in the 15.0-150  keV band.  For HMXBs this energy range is  extremely  important  as  resonant cyclotron  scattering  features (CRSFs)  are usually  observed  in  this part  of  the X-ray  spectrum \\citep{heindl04}.  In this  paper we analyse  the soft and  hard X-ray data  collected by BAT,  INTEGRAL and  by pointed  Suzaku and  Swift-XRT  observations on \\src.  This  source was discovered by INTEGRAL  in 2004 \\citep{bird04} and it was initially  associated with the radio pulsar PSR~J1649--4349 because of a spatial coincidence.  A later INTEGRAL observation with a deep  exposure allowed  to reduce  the positional  uncertainty  and to reject   the   pulsar   association  \\citep{atel457}.    A   follow-up observation with \\chandra found a soft X-ray counterpart at a position of RA(J2000)  = 16\\degr 49\\arcmin 26.92\\arcsec;  Dec(J2000) = -43\\degr 49\\arcmin  8.96\\arcsec \\citep{atel654}  allowing the  association with the   infrared   counterpart   2MASS~J1642695--4349090,  a   B0.5   Ib supergiant.   \\citet{nespoli10} performed  K-band spectroscopy  of the inferred counterpart,  confirming the B0.5\u20131  Ia-Ib companion spectral type; the infrared extinction, translated into the equivalent hydrogen column  density,  was  estimated  to  be (2.92  $\\pm$  1.96)  $\\times$ 10$^{22}$ \\cmmdue,  a value  that is lower  with respect to  the X-ray absorption, thus indicating  that part of the X-ray  absorption may be local to  the compact object.  \\citet{cusumano10}  found a periodicity of  6.732 $\\pm$  0.002 d,  interpreted as  the orbital  period  of the system. The folded  light curve shows presence of  an eclipse, lasting $\\sim$ 12\\% of the P$_{orb}$. Assuming a 1.4 \\msun\\, neutron star (NS) and 32 \\rsun\\, for the companion  star, the system has most likely a low eccentricity value ($e  \\leq$ 0.15).  Recently, \\citet{atel2766} found in RXTE observations evidence for a shorter periodicity at 1069 $\\pm$ 7 s, interpreted as the NS spin  period.  A spectral analysis using non simultaneous data from Swift-XRT and INTEGRAL \\citep{hill08} showed that the broadband X-ray spectrum could be well fitted  by an  absorbed (\\nh=5.6  $\\times$ 10$^{22}$  \\cmmdue) power-law ($\\Gamma$ = 0.6$\\pm$0.3) with a  high-energy cut-off at $\\sim$ 17 keV. \\citet{morris09}   analysing  Suzaku   data  of   \\src,   proposed  an alternative modelling of the  X-ray spectrum using a partial covering component multiplied by a  simple power-law. The covering fraction of the total  emission was estimated  in 0.62$\\pm$0.07, while  the local value of the \\nh was estimated in $\\sim$ 30 $\\times$ 10$^{22}$ \\cmmdue. ",
        "conclusions": "}  There are 15 sources that clearly have shown CRSFs signatures in their spectra. This number represents about  10\\% of the total population of HMXBs  detected   in  our  Galaxy  \\citep{reig11}   and,  despite  the considerable  amount of observing  time dedicated  by modern  and past X-ray observatories, to this class of sources, new detections of CRSFs have not considerably risen in the last decade.  We  have presented  new  results based  on  the analysis  of the  data collected by  BAT during the first 54  months of the \\sw  mission on a persistent  sgHMXB  \\src.   We  have  shown  that  the  long-term  BAT collected  hard X-ray  spectra provides  an access  for  detecting the presence of CRSFs in this kind of sources. This new approach relies on a  series  of  assumptions   and  caveats  that  will  be  hereinafter discussed.  The  broadband  spectrum, modelled  with  a positive-negative  cut-off power-law model  provides the statistically  most favoured description of the data in the examined datasets.  The residuals in the BAT energy range show the presence of  an absorption feature between $\\sim$28 and $\\sim$38 keV.   The feature is  not sensibly dependent on  the adopted soft X-ray contribution, although we  have just exploited the only two available to date pointed observations with XRT and Suzaku.  The broad negative residuals are compatible with a resonant cyclotron absorption feature with  an estimated significance  of the detection at  least at $\\sim$ 4 $\\sigma$ level.  Similar  features are  often  seen  in the  high  energy spectra  of sgHMXB.  They  are interpreted as cyclotron lines  produced near the magnetic poles of the  accretion-powered neutron star. These are due to resonant  scattering processes of  the X-rays by  electrons whose kinetic  energies are  quantised  in discrete  Landau energy  levels perpendicular to the B-field. Their detection is fundamental for the understanding of  the accretion mechanisms unto the  neutron star as it allows a direct measurement  of the magnetic field of the neutron star (or at  least a lower limit, depending on  the height where the resonant    line   absorption    is   effectively    produced,   see \\citealt{nishimura11}).  We  consider here  that the  cyclotron absorption  takes place  at the poles  of  the neutron  star,  taking  into  account the  relativistic gravitational  redshift: $\\rm  E_{\\rm  cyc}^{\\rm obs}$  = $\\rm  E_{\\rm cyc}(1 + \\rm z)^{-1}$, with  \\begin{equation} (1 + z)^{-1} = \\left( 1- \\frac{2GM_{\\rm X}}{R_{\\rm X}c^2}\\right)^{0.5} \\end{equation}  Assuming typical values for the NS  mass and radius (1.4 \\msun and 10 km, respectively) and the observed  value of the cyclotron line energy of 33$\\pm$4  keV (\\texttt{npex}  model for \\textit{Spectrum  1}), this implies a magnetic field of B$_{\\rm NS}$ (3.7$\\pm$0.4) $\\times10^{12}$ Gauss at the surface of the NS.  The   energy  resolution  of   the  spectra   is  not   sufficient  to independently constrain  line energy,  line width and  depth.  Because the line energy and depth  could be relatively better constrained with respect to the  line width, we choose to freeze this  value to 10 keV. This value represents a  first-order guess assumed in consideration of the  emerging relations  among  the parameters  that characterize  the cyclotron line  shapes.  An empirical correlation  between line energy and line width \\citep{coburn02} is, in fact, observed, where cyclotron lines detected between  30 keV and 40 keV  show preferably line widths in  the  5--10 keV  range  (e.g.  in  MX  0656-072,  4U 0352+309,  XTE J1946+274,  [see \\citet{coburn02,mcbride06}]).   We tried  in  all the fittings also to freeze the line width at 5 keV and 8 keV to check how sensible was  the determination of this parameter.  We generally found that the 10 keV guess always provided the lowest $\\chi^2$, even if the improvements  in  most  cases  were  marginal  (2$<\\Delta  \\chi^2<$5). Another important empirical  correlation concerns the cut-off energies of  the continuum  emission and  the  cyclotron line  energy, where  a correlation is  observed with cut-off  energies typically at  half the value of the cyclotron energy.   Also in this case, our results appear in agreement with  the general trend that is  observed, with a cut-off energy at  $\\sim$15 keV and the  cyclotron energy at  $\\sim$ two times this value \\citep{heindl04}.  Possible biases in this analysis are envisaged by the use of long-term time  averaged spectra for  features that  may show  some variability. Some luminosity-dependent  shifts of the cyclotron  line position were in      fact      clearly      observed      in      some      sources \\citep{mihara04,staubert07,nakajima10}.      However     the     large uncertainties in  the present analysis  estimates do not allow  to set this  possibility into  investigation.  \\src  has never  shown unusual periods  of strong  changes  in its  luminosity,  nor any  outbursting behavior as  in the  case of X0331+53,  where luminosity changed  by a factor $\\sim$  200 \\citep{nakajima10}, so that we  argue that possible shifts,  if   present,  should   be  within  our   quoted  error-bars. Moreover, it is  to be noted that the  cyclotron feature at 45 keV of  A0535+262 did  not show any  variations despite a  change in luminosity  of two orders  of magnitude  \\citep{terada06}. Another source of bias concerns the possibility of continuum spectral changes, that, when averaged, could  result in spectral artifacts. The analysis of  the  soft   X-ray  spectrum  from  the  XRT   and  Suzaku  pointed observations are not sufficient to test any variability pattern in the broadband  spectral properties;  use of  simplified, phenomenological, models are  also possible source of  bias and a  more physically based model should be employed to  also infer the physical properties of the accretion  column   environment  \\citep{becker07,schonherr07}  and  to better constrain the shape of cyclotron resonance features.   \\subsection{Conclusions} We have presented  spectral analysis in the high-energy  X-ray band of \\src using long-term survey data from BAT and ISGRI, complemented with pointed soft X-ray observations from XRT and Suzaku.  The BAT spectrum is to date the highest S/N  spectrum of this source in the high-energy X-ray band.  When the spectrum is complemented with soft X-ray data, a broad absorption  feature is  detected, irrespective of  the broadband model used to  fit the data.  We modeled the  feature assuming that it is a  CRSF. This interpretations  is statistically acceptable  and the physical interpretation  of the parameters  is in agreement  with what expected from the whole sample of HMXBs that show similar features. We estimated  the significance level  of the  detection $>$  3.9 $\\sigma$ confidence level and next generation of future observatories dedicated to the high X-ray band  (e.g.  NuSTAR and Astro-H) will likely provide a more stringent confirm of its presence and a substantial improvement in the determination of its shape."
    },
    "1208/1208.0443_arXiv.txt": {
        "abstract": "We present a parametrized study of the effects of free thermal neutron injection on primordial nucleosynthesis, where both the rate and the time scale of injection are varied. This generic approach is found to yield a successful solution for reducing the \\sep\\ abundance without causing significant problems to other elemental abundances. Our analysis demonstrates that hadronic injection, possibly due to decays or annihilations of dark matter particles with a mass of about 1 to 30 GeV, provides a possible solution to an outstanding problem in the standard Big Bang model. ",
        "introduction": "The motivation for this study concerns the discrepancy between the primordial Li abundance predicted in the canonical Big Bang model and observational data. The primordial lithium abundance is deduced from observations of low metallicity stars in the halo of our Galaxy where the lithium abundance is almost independent of metallicity, displaying a plateau, the so-called Spite plateau \\citep{Spite82}. This interpretation assumes that lithium has not been depleted at the surface of these stars, so that the presently observed abundance is supposed to be equal to the initial value. The small scatter of values around the Spite plateau is an indication that depletion may not have been very effective. Astronomical observations of these metal-poor halo stars \\citep{Ryanetal2000} have led to a relative primordial abundance of:  Li/H~$ = (1.23^{+0.34}_{-0.16}) \\times 10^{-10}$.  \\noindent A more recent analysis by \\citet{sbordone10} gives:  Li/H~$ = (1.58 \\pm 0.31) \\times 10^{-10}$.  \\noindent More generally, \\citet{Spite10} have reviewed recent Li observations and their different astrophysical aspects. Also see \\citet{frebel11} for a comprehensive review.  On the other hand, the most recent \\sbbn\\ (SBBN) calculations, using the most up-to-date nuclear data, give:  Li/H = $(5.14\\pm0.50)\\times 10^{-10}$ \\citep{CV10}.  \\noindent Hence there is a factor of 3-4 discrepancy between observation and theory at the WMAP7 baryonic density.  \\sep\\ is produced as a by-product of decay of $^7$Be. Nuclear mechanisms to destroy this $^7$Be have been explored. \\addCorr{A possibly} increased $^7$Be(d,p)$2\\alpha$ cross section has been proposed by \\citet{Coc04} and later by \\citet{Cyb09} but was not confirmed by experiments \\citep{Ang05,OMa11,Kir11}. Other $^7$Be destruction channels have recently been proposed by \\citet{Cha11} and await experimental investigation.  Another scenario would be to take advantage of an increased late-time neutron abundance, as introduced in~\\cite{Reno:1987qw} for the generic case of hadronic injection. In the context of varying constants, when the $^1$H(n,$\\gamma)^2$H rate is decreased, the neutron late-time abundance is increased (with no effect on \\qua) so that more $^7$Be is destroyed by $^7$Be(n,p)$^7$Li(p,$\\alpha)\\alpha$, \\citep[see in] [Fig. 1]{Coc07}. Many other nuclear reactions could be potential sources of free neutrons. However, a recent study \\cite{Coc11} extended the SBBN network to 59 nuclides from neutrons to $^{23}$Na, linked by 391 reactions involving n, p, d, t and \\tro\\ induced reactions and 33 $\\beta$-decay processes. The \\sep\\ abundance is now estimated to Li/H = $5.24\\times 10^{-10}$ \\citep{Coc11}, as found also by \\citep{Cyb08}. This confirms and even increases the discrepancy.  Including physics beyond the standard model of particle physics and beyond the standard Big-Bang picture can also give rise to extra neutron injection. Indeed, BBN can be used as an anchor to test the plausibility of new physics, and conversely, new physics can provide mechanisms to help solving the SBBN discrepancies with observations~\\cite{Malaney:1993ah,Sarkar:1995dd,Jedamzik:2009uy,Pospelov:2010hj}. One such option is that of hadronic decays of exotic unstable particles. For example, a metastable stop Next-to-Lightest Supersymmetric Particle (NLSP) decays into a gravitino Lightest Supersymmetric Particle (LSP), thus a dark matter candidate, and a top quark injects energetic protons and neutrons during nucleosynthesis~\\cite{PhysRevD.79.043514,1988ApJ...330..545D,Reno:1987qw,Jedamzik:2004er,Jedamzik06,Cumberbatch:2007me,Cyburt:2010vz,Pospelov:2010cw}. Another possible source of neutrons arises from residual annihilations of dark matter particles -- such as neutralino LSP annihilating into fermion-antifermion couples that further hadronize -- that are chemically decoupled at BBN times~\\cite{Reno:1987qw,Jedamzik:2004er,Pospelov:2010hj}. In all these scenarios, neutron injection provides the primary impact on BBN and Li production.  \\addCorr { In this work, we study the effect of free neutron injection, parametrized by the injection rate and time-scale. Different injection models are thus included in the full code presented in~\\cite{Coc11}. Hence, this implementation is expected to give hints regarding the injection mechanism including possible nuclear reaction uncertainties, fundamental constant variations and exotic particle decays or annihilations. We comment on possible scenarios behind neutron injection, however, we do not include a full treatment of the production and thermalization of neutrons in the code.} ",
        "conclusions": "While in our SBBN code, the neutrons are injected at equilibrium, it is likely that extra neutrons from any kind of beyond the Standard Model physics are produced out-of-equilibrium. It is, therefore, important to consider the thermalization process of neutrons during BBN before they decay, through which channels they do so, and to estimate the possible perturbations to SBBN abundances. As mentioned by \\citet{Jedamzik:2004er,Jedamzik06}, the thermalization process calculation of energetic nucleons is simplified greatly by two facts: first, the Hubble time is much greater than the mean time between any of the interactions under consideration $\\tau_H \\approx 300 \\left(\\frac{T}{90 \\, \\mathrm{KeV}}\\right)^{-2} \\mathrm{s}$, and second, the interactions between non-thermal and thermal nucleons are unlikely. The decay time of free neutrons ($\\tau_0 = 881 \\,\\mathrm{s}$) is even greater than the Hubble time.  \\addCorr{There are} 3 main classes of reactions: (1) elastic and inelastic $\\mathrm{n}-\\mathrm{p}$ scattering, (2) the afore-mentioned spallation of ${}^4\\mathrm{He}$ with production of ${}^3\\mathrm{He}$, and (3) both elastic and inelastic scattering $\\mathrm{n}-{}^4\\mathrm{He}$. All of these processes contribute to thermalization, but the spallation to non-thermal ${}^3\\mathrm{He}$ might disturb the abundance of ${}^6\\mathrm{Li}$ \\cite{Ellis11} through the following reactions \\begin{equation} \\begin{array}{lllllll} n &+& {}^4\\mathrm{He} &\\to& {}^3\\mathrm{He} &+& 2n \\\\ {}^4\\mathrm{He} &+& {}^3\\mathrm{He} &\\to& {}^6\\mathrm{Li} &+& p. \\end{array} \\end{equation}  Firstly, we justify the claim that the injected neutrons indeed thermalize before they decay and secondly, we estimate the production of ${}^3\\mathrm{He}$ and ${}^6\\mathrm{Li}$.  \\addCorr{\\six\\ is a very interesting isotope as a new cosmological nucleus. Indeed, its abundance is measured in low metallicity stars and offers a unique probe of two different mechanisms of nucleosynthesis: SBBN and cosmic rays. The former producing predominantly \\sep, \\six\\ was until recently considered as a pure spallative --i.e., post BBN-- product. However, according to recent detections in very low metallicity PopII stars, the average is \\six/\\sep$\\simeq 0.042$~\\cite{Asplund:2005yt}. These observations have been interpreted as evidence for a large primitive abundance of \\six\\ (\\six/H~$\\simeq 10^{-12}$) while SBBN calculations confirm a low primordial value (\\six/H~$\\simeq 10^{-14}$). For details on the subject, see~\\cite{Coc11,PhysRevC.82.065803}. Recently, new studies using 3D atmosphere model in metal poor halo stars reconsider the detection of \\six. In~\\cite{Steffen2012} two detections are confirmed (\\six/\\sep\\ $\\simeq5-10$\\%). More observations are presently needed to improve the statistics. Nevertheless, this new spectroscopic research can be an indicator of new physics, as has been point out by~\\citet{Jedamzik:1999di}. In that regard, we have to make sure that the extra physics we consider does not perturb significantly the abundance of \\six. We do it in an order of magnitude estimate. Figure~\\ref{f:cno} shows clearly that our models should not modify too much \\six: at $\\lambda_0\\simeq10^{-8}$~n~s$^{-1}$ \\six\\ is in the range $(1-6)\\times 10^{-14}$.}  The total cross-section of elastic and inelastic scatterings of n off p is about 70 mb at 100 MeV and 30-40 mb above 100 MeV. The cross-section of the reaction $\\mathrm{n} + {}^4\\mathrm{He} \\to {}^3\\mathrm{He} +2\\mathrm{n}$ varies from about 15 mb at 30 MeV and 50 mb at 50 MeV down to 20 mb at 1 GeV. The sum of the cross-sections of the remaining inelastic processes is comparable to $\\mathrm{n} + {}^4\\mathrm{He} \\to {}^3\\mathrm{He} +2\\mathrm{n}$ at 50 MeV and about 3-4 times greater than the rates above 100 MeV. The elastic scattering of neutrons off ${}^4\\mathrm{He}$ varies from 500 mb at 20 MeV down to 30 mb at 100 MeV and stays constant at higher energies.  The characteristic time of scattering of a neutron of kinetic energy $E_k$ off thermal H and ${}^4\\mathrm{He}$ is \\begin{align} \\tau_\\sigma \\approx A_i \\!\\left(\\frac{T}{90\\,\\mathrm{KeV}}\\right)^{-3}\\! \\left(\\frac{\\sigma}{50\\,\\mathrm{mb}}\\right)^{-1} \\nonumber \\!\\left(\\frac{1\\!+\\!(m_n/50\\,\\mathrm{MeV})^2}{1+(m_n/E_k)^2}\\right)^{-1/2} \\mathrm{s}, \\end{align}  where $A_H = 26$ and $A_{^4He} = 2.17$, $E_k$ is the kinetic energy of the neutron. In the relativistic limit, $E_n \\gg m_n$, $\\tau$ goes down to 0.07 s, while the decay time grows linearly with energy $\\tau_n = \\frac{E}{m_n}\\tau_0$.  We also have to make sure that a significant fraction of energy is transferred in a single scattering or per mean free path length (its inverse is denoted $E (\\lambda dE/dx)^{-1}$ in \\citet{Jedamzik06}). In inelastic $\\mathrm{n}-\\mathrm{p}$ scatterings, $\\frac{\\Delta E}{E}$ remains constant up to 250 MeV ($\\frac{\\Delta E}{E} \\gtrsim 0.1$). Indeed in elastic $\\mathrm{n}-{}^4\\mathrm{He}$ scatterings, the average energy transfer is about $5-10$ MeV. In inelastic processes $\\mathrm{n}-{}^4\\mathrm{He}$, the neutron loses as much as the binding energy of ${}^4\\mathrm{He}$ (28.3 MeV) and a quarter of the remaining energy.   From these arguments, it is clear that extra neutrons thermalize in these conditions before decaying in a range of temperatures from 100 KeV down to a few KeV and in the range of energies from about 10 MeV to 1 GeV.  The hypothesised extra neutrons might be produced by annihilating dark matter particles. The energy injection due to dark matter annihilations if the freeze-out of the dark matter species happened at BBN temperature is severely constrained (see~\\cite{Sarkar:1995dd} for example). However, if the freeze-out happened before BBN, annihilations become marginal, as the expansion rate dominates the interaction rate. Nonetheless there would be a residual annihilation rate of dark matter into standard model particles. Eventually, after hadronization, a spectrum of neutrons would be generated, that would reach thermal equilibrium as discussed before.  The annihilation rate of uniformly distributed dark matter per baryon can be written as \\begin{equation} \\begin{array}{lll} \\Gamma_b &=&  \\frac{1}{2}\\langle \\sigma v \\rangle \\frac{n^2_{0,X}}{n_{0,b}} \\left(1+z\\right)^3 \\\\ &=& 5.3\\times 10^{-9}\\left(\\frac{\\langle \\sigma v \\rangle}{3\\times10^{-26} cm^3 s^{-1}}\\right) \\left(\\frac{30 \\, \\mathrm{GeV}}{M_X}\\right)^2 \\left(\\frac{T}{90\\,\\mathrm{keV}}\\right)^3 s^{-1}, \\end{array} \\end{equation}  where $n_{0,X}$ and $n_{0,b}$ are the present day number densities of dark matter and baryons. We see that at a temperature of about $90\\;\\mathrm{keV}$, a particle dark matter mass of $M_X \\la 30 \\,\\mathrm{GeV}$ and a canonical annihilation rate are plausible parameters needed to achieve $\\lambda = (3-5) \\times 10^{-9}$~s$^{-1}$, depending on the neutron spectrum generated by the annihilations.  The neutron spectrum is generated after hadronization of the particles produced at annihilation, and it is expected to be peaked at roughly $M_X/5 - M_X/15$. Therefore, the lighter the dark matter particle, the larger the fraction of thermal neutrons. However, the dark matter has to be heavy enough to produce neutrons, hence, the most interesting mass range lies roughly between 1 and 30 GeV.  The annihilation rate and branching ratios depend on the dark matter candidate. Moreover, the dark matter temperature evolves from chemical decoupling down to thermal decoupling (see~\\cite{Bringmann:2006mu}). The dependence of the annihilation rate on the dark matter temperature can be very strong; for example, if the freeze-out mechanism invokes a nearly resonant exchange, or co-annihilations~\\cite{Griest:1990kh}.  A relevant example for a dark matter candidate in the mass range discussed here is the neutralino in the Next-to-Minimal Supersymmetric Standard Model. As shown in~\\cite{AlbornozVasquez:2011js}, the resonant mechanism at freeze-out can yield a very large boost to the annihilation rate at lower temperatures (see Fig. 4 in~\\cite{AlbornozVasquez:2011js}). For kinetic decoupling at $T_{kd}\\sim T_{fo}/10$, one could have a factor $10-100$ enhancement in the annihilation rate from the $3\\times10^{-26}\\rm{cm}^3\\rm{s}^{-1}$ required at freeze-out. Thus there can be a variety of dark matter candidates (with different masses and annihilation cross-section mechanisms) which provide an injection flux of $(3-5) \\times 10^{-9}$~s$^{-1}$.  \\addCorr{It is interesting to note that some of these candidates could explain direct detection signals as they have the right mass range and could attain the needed interaction rates with detectors. Also, they could be challenged by $\\gamma$-ray production at dwarf spheroidal galaxies. Relating annihilating rates at freeze-out, BBN and galactic times, and elastic scattering interactions with nuclei, can provide powerful constraints on a given dark matter model.}   In conclusion, neutron injection can help to resolve the \\sep\\ problem provided that the neutrons are essentially thermal. This can be achieved for annihilations or decays of dark matter particles in the mass range 1-30 GeV. A detailed physical model involving, for example, a metastable supersymmetric NLSP or annihilating neutralino dark matter is beyond the scope of this paper, but would seem to be easily achievable."
    },
    "1208/1208.0733_arXiv.txt": {
        "abstract": "In continuation of their earlier measurements, the PAMELA group reported data on antiproton flux and $\\overline{P}/P$ ratios in 2010 at much higher energies. In past we had dealt with these specific aspects of PAMELA data in great detail and each time we captured the contemporary data-trends quite successfully with the help of a multiple production model of secondary antiprotons with some non-standard ilk and with some other absolutely standard assumptions and approximations. In this work we aim at presenting a comprehensive and valid description of all the available data on antiproton flux and the nature of $\\overline{P}/P$ ratios at the highest energies reported so far by the PAMELA experiment in 2010. The main physical implication of all this would, in the end, be highlighted. \\bigskip \\par Keywords: Cosmic ray interactions. Composition, energy spectra and interactions. Cosmic rays (including sources, origin, acceleration, and interactions). Dark Matter (stellar, interstellar, galactic, and cosmological). \\\\ \\par PACS nos.: 13.85.Tp, 98.70.Sa, 95.35.+d ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.4110_arXiv.txt": {
        "abstract": "The spatial-frequency coverage of a radio interferometer is increased by combining samples acquired at different times and observing frequencies.  However, astrophysical sources often contain complicated spatial structure that varies within the time-range of an observation, or the bandwidth of the receiver being used, or both.  Image reconstruction algorithms can been designed to model time and frequency variability in addition to the average intensity distribution, and provide an improvement over traditional methods that ignore all variability. This paper describes an algorithm designed for such structures, and evaluates it in the context of reconstructing three-dimensional time-varying structures in the solar corona from radio interferometric measurements between 5 GHz and 15 GHz using existing telescopes such as the EVLA and at angular resolutions better than that allowed by traditional multi-frequency analysis algorithms. ",
        "introduction": "\\label{sec:intro}  %  Image reconstruction in radio interferometry\\cite{NRAO_LECTURES} benefits from maximizing the number of distinct spatial frequencies at which the visibility function of the target is sampled.  For a given array, the spatial-frequency coverage can be increased by taking measurements across several hours (Earth-rotation-synthesis) and at multiple observing frequencies (multi-frequency-synthesis).  However, astrophysical sources sometimes show time and frequency variability within the parameters of such an observation.  Reconstruction algorithms have traditionally ignored this variability, leading either to imaging-artifacts or to data-analysis strategies that use subsets of the data independently, both which limit imaging sensitivity and accuracy.  In recent years, broad-band receivers have been installed on radio interferometers specifically to increase imaging sensitivity, and this has necessitated the development of algorithms\\cite{MFCLEAN_CCW, MFCLEAN, MSMFS2011} that account for sky spectra in addition to average intensity. A similar idea applies to time-variability as well, and recent work\\cite{MBCLEAN2011} has demonstrated combined time and frequency modeling for point-sources (without multi-scale support), mainly for the purpose of eliminating artifacts from images of the average intensity distribution.  This paper presents a reworking of the system of equations being solved when modeling time and frequency varying structure, and includes multi-scale support. Using a software implementation of a multi-scale, multi-frequency, time-variable image-reconstruction algorithm that can be operated in modes that emulate various intermediate algorithms, a series of imaging tests were carried out on data simulated for three-dimensional time-varying structure in the solar corona. The results are discussed in terms of the accuracy of reconstructing the time and frequency-dependent structure at an angular resolution better than what traditional methods allow. ",
        "conclusions": "Image-reconstruction algorithms for radio interferometry that model time and frequency structure can be used not-only to improve the fidelity of the average image, but to also reconstruct the frequency-dependent and time-variable structure at an accuracy sufficient for astrophysical interpretation.   Smooth frequency-dependent structure can be reconstructed across frequency ranges wider than previously demonstrated, achieving higher angular-resolution than that given by the lowest frequency in the measured band.  Reconstructing time-variability is also possible, but since angular-resolution is not of concern here, this feature may be useful only in the context of improving signal-to-noise and image-fidelity by using all measurements together.  The example in Sec.~\\ref{sec:sunspot} is a proof-of-concept demonstration that such imaging of frequency-dependent and time-variable structures in coronal active regions may be possible with current telescopes such as the EVLA.  Future telescopes such as FASR will have more than enough spatial-frequency coverage for an unambiguous reconstruction at each timestep and channel, even after tapering data from all frequencies to a common low angular resolution.  From a practical standpoint, the MS-MF-TV-CLEAN implementation described in this paper improves upon the multi-beam algorithm in Ref.~\\citenum{MBCLEAN2011} in that it does not suffer from instabilities incurred when there is significant overlap between sampled spatial-frequencies. It suggests the choice of $N_p + N_q - 1$ basis functions instead of the $N_p \\times N_q$, and also folds-in multi-scale support in a way that allows a wide choice of spatial basis functions whose amplitudes are given by time and frequency polynomials.  There is always a trade-off between a basis set that optimally represents the structure being reconstructed and requires a minimal number of components to model it, and a basis set whose members the instrument can optimally distinguish between (an orthogonal basis for the instrument) but may require many more components to model the target structure.  For the situations explored in this paper, very simple polynomial basis functions produce usable reconstructions, although some care must be taken to ensure that the instrument can distinguish between them. This information is encoded in the diagonals of each block of the Hessian matrix, and scale-sizes and polynomial orders must be chosen such that its condition number is not too large. Fine structures in the dynamic spectra can be reconstructed only with higher-order terms which are increasingly harder for the instrument to distinguish from one another, and as demonstrated in Ref.~\\citenum{MBCLEAN2011} may require an explicit orthogonalization step. However, this orthogonalization must be restricted to only the time and freqency basis functions because the multi-scale basis functions are usually chosen based on structures that optimally describe the structure being imaged, and a basis function orthogonalization may result in a much larger number of components being required to fit the data.  One problem with this approach is the current inability to derive accurate error-estimates on the fitted parameters, because it depends not-only on the basis functions and the instrument (as encoded in the Hessian and covariance matrices), but also on the appropriateness of the chosen model for the structure being reconstructed. For instance with MS-MF-CLEAN\\cite{MSMFS2011} , errors in the spectral-index data product are as high as a few hundred percent when delta-functions are used to model extended emission.  There is also an inherent ambiguity in the spectral structure at the largest scales, because of the frequency-scaling of the instruments $uv$-coverage, and in this case, the errors are dominated by the inability of the measurements to constrain the model and may require {\\it a-priori} information\\cite{MSMFS2011, FASRImaging2003} . To be practically usable, better measures of the reconstruction quality are required.  The ideas presented in this paper are one step closer to the direct modeling goals presented in Ref.~\\citenum{Fleishman2010Concept} for the reconstruction of 3D magnetic-field structures in the solar corona . The next step for such algorithms ought to be to step away from pattern-modeling, and fold physical models directly into the reconstruction process. Refs.~\\citenum{Fleishman2009} and \\citenum{Fleishman_modeling2011} , among others, demonstrate that physical models fitted to single-channel images from an idealized interferometer are capable of recovering magnetic-field structures and other information required to undestand the physical processes. An algorithm that combines these ideas and reconstructs the physical model directly from visibilities would eliminate the intermediate non-linear step of imaging which has its own uncertainties.   \\subsection"
    },
    "1208/1208.4260_arXiv.txt": {
        "abstract": "Using GMOS--IFU spectroscopic observations of the compact H\\,{\\sc ii}/BCD galaxies Tol 0104-388 and Tol 2146-391, we study the spatial distribution of emission lines, equivalent width EW(H$\\beta$), extinction c(H$\\beta$), ionization ratios ([OIII] $\\lambda$5007/H$\\beta$, [SII] $\\lambda\\lambda$6717,6731/H$\\alpha$ and [NII] $\\lambda$6584/H$\\alpha$), kinematics, and the chemical pattern (O/H, N/H and N/O) of the warm interstellar medium in these galaxies. We also investigate a possible dependence of these properties on the I(He \\,{\\sc ii} $\\lambda$4686)/I(H$\\beta$) ratio and find no significant correlation between these variables. In fact, the oxygen abundances appear to be uniform in the regions where the He \\,{\\sc ii} $\\lambda$4686 emission line was measured. It can be interpreted in the sense that these correlations are related to global properties of the galaxies and not with small patches of the interstellar medium. Although a possible weak N/H gradient is observed in Tol 2146-391, the available data suggest that the metals from previous star-formation events are well mixed and homogeneously distributed through the optical extent of these galaxies. The spatial constancy of the N/O ratio might be attributed to efficient transport and mixing of metals by starburst-driven super-shells, powered by a plethora of unresolved star cluster in the inner part of the galaxies. This scenario agrees with the idea that most of the observed He \\,{\\sc ii} $\\lambda$4686 emission line, in our sample of galaxies, is produced by radiative shocks. ",
        "introduction": "The ionizing radiation from newly formed stars and its interaction with the surrounding gas generate collisionally excited and recombination emission lines, that are commonly observed in starburst and low metallicity \\citep[1/50Z$_{\\odot}$-1/3Z$_{\\odot}$;][]{KS83} dwarf galaxies, such as, H\\,{\\sc ii} or blue compact dwarf (BCD) galaxies. From recent studies it was concluded that the hardness of the ionizing radiation from the current star formation (SF) activity increases with decreasing metallicity \\citep[e.g.,][]{G00,S03,TI05} as is expected in primeval galaxies in the early universe. These first stars (population \\,{\\sc iii} stars) should be very massive and hot \\citep[e.g.,][]{A02}, emitting very hard ionizing radiation, thus very effective in ionizing hydrogen and helium, then strong He \\,{\\sc ii} emission lines are likely present in the spectra of these galaxies \\citep[e.g.,][]{S02,S03}. Originally, the low abundances of heavy elements in HII/BCD galaxies and the non-detection of an old stellar population have given rise to the question of whether they may be presently forming their first generation of stars \\citep{SS70}. Recent works, however, have shown that most H\\,{\\sc ii} galaxies seem to present an underlying population of old stars from previous episodes of SF \\citep[e.g.,][]{P96,TT97,C03} suggesting an intermittent SF history with short intense SF episodes followed by long quiescent phases.   \\begin{figure} \\includegraphics[width=70mm]{fig1.eps} \\caption{g-band acquisition image of the galaxies Tol 0104-388 and Tol 2146-6391 in logarithmic scale. The rectangle indicates the GMOS--IFU FoV of 3\\arcsec.5$\\times$5\\arcsec. } \\label{image} \\end{figure}  To date, a small but an increasing number of H\\,{\\sc ii}/BCD galaxies has been observed using Integral Field Unit (IFU) spectroscopy \\cite[e.g.,][]{I06a,L09,C09a,C09b,C10,MI10} in order to study the spatial distribution of their properties (i.e., emission lines, extinction, kinematics, abundances). In particular, some of these studies have shown a remarkable chemical homogeneity of oxygen abundance in BCD galaxies \\citep[e.g.,][]{L09}. This implies that oxygen and all $\\alpha$-elements are primary, produced by massive stars ($>$8-10 M$_{\\odot}$) and released into the interstellar medium (ISM) during their supernova (SN) phase. These newly synthesized metals will be dispersed in the whole galaxy and mixed via hydrodynamic mechanisms in timescales of few 10$^{8}$ yr. On the other hand, the ratio N/O was also found to be rather constant for BCD galaxies \\citep[log(N/O)$\\simeq$-1.6;][]{EP78,A79,I99}, implying mainly primary production of nitrogen at low metallicity (12+log(O/H)$\\leq$7.6), although the amount coming from each source, primary or secondary, is still debated because of the lack of a clear mechanism that produces N in massive stars besides the effect of the stellar rotation \\citep{MM05}. The observed spread of N/O at 7.6$\\leq$12+log(O/H)$\\leq$8.3 is seen to be large and has been attributed to observational uncertainties, a loss of heavy elements via galactic winds \\citep{V98} and/or to the time delays (delayed-release hypothesis) between the production of oxygen by massive stars and that of nitrogen by intermediate and/or massive stars. In this scenario the SF is an intermittent process in H\\,{\\sc ii}/BCD galaxies \\citep{G90} with several SF bursts. However, the delayed-release scenario cannot explain the presence of H\\,{\\sc ii}/BCD galaxies with high N/O ratio in comparison to the expected value for their O content. The most plausible explanation of the high N/O ratio observed in these objects is the chemical pollution of N due to the presence of Wolf-Rayet (WR) stars \\citep[e.g.,][]{LS10a}. Localized nitrogen self-enrichment has been measured in a few cases \\citep[e.g.,][in NGC 5253]{K97} and attributed to the release of N into the ISM by the action of strong winds produced by these stars. Although, there are a few BCD galaxies with significant variations of oxygen abundance, such as in two HII regions of Haro 11 \\citep[12+logO/H=8.33$\\pm$0.01 in Haro B and 12+logO/H=8.10$\\pm$0.04 in Haro C;][]{G12}, no localized oxygen enrichment systems have been confirmed \\citep[][]{L09}. Finally, at high metallicity (12+log(O/H)$\\geq$8.3) the N/O ratio clearly increases with increasing oxygen abundance. Hence, nitrogen is essentially a secondary element in this metallicity regime.  The origin of hard-ionization emission lines such as He \\,{\\sc ii} $\\lambda$4686 in BCD galaxies has been a subject of study in recent years given that photoionized models of H\\,{\\sc ii} regions fail to reproduce the observed high-ionization line ratios. In particular, the observed intensity of He \\,{\\sc ii} $\\lambda$4686 with respect to H$\\beta$ is several orders of magnitude larger than model predictions for photoionized regions \\citep{S90}. Several mechanisms for producing hard ionizing radiation have been proposed, such as massive main-sequence stars \\citep{S97}, WR stars \\citep{S96}, primordial zero-metallicity stars \\citep[][]{S02,S03}, high-mass X-ray binaries \\citep[HMXBs;][]{G91}, radiative shocks \\citep{DS96} and O stars at low metallicity may also contribute to the He \\,{\\sc ii} ionizing flux \\citep{B08}. \\citet{F01} and \\citet[][]{I01,I04} have explored these different mechanisms which can produce the hard radiation in SBS 0335-052, Tol 1214-277 and Tol 65. They concluded that the ionization produced by main-sequence stars cannot explain the strong [Fe V] $\\lambda $4227 and He \\,{\\sc ii} $\\lambda $4686 emission lines, but other ionization sources as WR stars, HMXB systems and fast shocks, can be considered.  We focus here on two compact H\\,{\\sc ii} galaxies: Tol 0104-388 and Tol 2146-391 in order to study the relationship between the physical-chemical and kinematics properties, the nature of their hard ionization radiation pattern (He \\,{\\sc ii} $\\lambda$4686 emission lines) in the ISM of these galaxies based on IFU spectroscopic observations on the Gemini South telescope. We selected these galaxies because both objects belong to a subset sample of H\\,{\\sc ii} galaxies with compact morphology and relative low metallicity (12+log(O/H)$\\la$7.8-7.9 dex), possibly mimicking the properties one expects for young galaxies at high redshift. In Fig.~1 we show g-band acquisition images of the galaxies Tol 0104-388 and Tol 2146-391. With the present observations, we note that Tol 2146-391 encompasses two main giant H\\,{\\sc ii} regions (GH\\,{\\sc ii}Rs), as shown in Fig.~1. While Tol 0104-388 shows one GH\\,{\\sc ii}R. In Table ~1 we show the general parameters of these galaxies.   \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{General parameters of the galaxies.} \\begin{tabular}{@{}lccccccrlr@{}} \\hline Name     &    \\multicolumn{2}{l}{Coordinates (J2000)}& D (3K CMB)\\footnote{Obtained from NED.} &  1\\arcsec (3K CMB) & 12+log(O/H)\\footnote{Derived from the present observations.} & Other names \\\\  \\hline Tol 0104-388 & 01:07:02.1 &  -38:31:52 & 88.4 & 429 &8.02 & CTS 1001 \\\\ Tol 2146-391 & 21:49:48.2 &  -38:54:09 & 117.3& 569 &7.82 & \\\\  \\hline \\end{tabular} \\end{minipage} \\end{table*}  This paper is organized as follows: the observations and data reduction are presented in Sect.~2. In Sect.~3 we show our results based in the study of the detected emission lines. In Sect.~4 we discuss our results and in Sect.~5 we summarize our conclusions. ",
        "conclusions": "Using new GMOS--IFU spectroscopic observations of the compact H\\,{\\sc ii}/BCD galaxies Tol 0104-388 and Tol 2146-391, we studied the spatial distribution of the high-ionization emission line He \\,{\\sc ii} $\\lambda$4686 and the chemical pattern through the ISM of the galaxies in an extended region of 3$\\arcsec$.2$\\times$5$\\arcsec$, equivalent to $\\sim$1372 pc $\\times$ 2058 pc and $\\sim$1820 pc $\\times$ 2730 pc for Tol 0104-388 and Tol2146-391, respectively. Based on the analysis of its properties, we have obtained the following results:  \\begin{enumerate} \\item The examination over each individual 0\\arcsec.2 spaxel and also in the integrated spectra, in both galaxies in our sample, do not reveal any clear stellar WR features.  \\item Both galaxies show the presence of the emission line He \\,{\\sc ii} $\\lambda$4686 with an integrated intensity relative to H$\\beta$ of I(He \\,{\\sc ii} $\\lambda$4686)/I(H$\\beta)<$0.02. We did not detect a clear correlation between the spatial distribution of EW(H$\\beta$), 12 + log (O/H) and log(N/O) with respect to the hardness of this high-ionization radiation across the spaxels in Tol 2146-391.  \\item Given the spatial distribution of He \\,{\\sc ii} $\\lambda$4686 emission in our two analyzed galaxies, this high ionizing radiation is likely associated with a mix of sources, where WR stars, HMXB and O stars cannot be excluded. While, expanding shells powered by a plethora of unresolved star clusters are likely producing most of the observed He \\,{\\sc ii} $\\lambda$4686 emission in our sample galaxies, at least in Tol 2146-391.  \\item We found some evidence that the 12+log(N/H) radial distribution, in Tol 2146-391, shows a slight trend, with the values decreasing with distance from H$\\alpha$ continuum peak. If real, this observed trend of 12+log(N/H) abundance would argue in favour that these heavy elements were produced during the previous burst of SF and are currently dispersed by the expansion in the ISM of starburst-driven super-shells. However, the spatial constancy of the N/O ratio might be attributed to efficient transport and mixing of metals by hydro-dynamical effects during the previous episodes of SF. \\end{enumerate}  All results presented here are suggestive that the physical conditions in these two galaxies, as in the case of the low luminosity and compact galaxy UM 408 \\citep{L09}, vary in a very small dynamical range and are quite homogeneus. Therefore, the lack of significant variation in abundance across the ISM of the galaxies would be indicative of uniform SF history occurring in galactic scales."
    },
    "1208/1208.3440_arXiv.txt": {
        "abstract": "We argue that global magnetic field reversals similar to those observed in the Milky Way occur quite frequently in mean-field galactic dynamo models that have relatively strong, random, seed magnetic fields that are localized in discrete regions. The number of reversals decreases to zero with reduction of the seed strength, efficiency of the galactic dynamo and size of the spots of the seed field. A systematic observational search for magnetic field reversals in a representative sample of spiral galaxies promises to give valuable information concerning seed magnetic fields and, in this way, to clarify the initial stages of galactic magnetic field evolution. \\bigskip  \\noindent {\\itshape Keywords:} Galactic dynamos; Magnetic fields of galaxies; Magnetic field reversals;  MHD ",
        "introduction": "\\label{intro} An outstanding question in the study of magnetic fields in spiral galaxies is that of large-scale field reversals. Do they occur, how easily can they be detected, can dynamo models explain them?  It appears that there may be such a field reversal in the Milky Way -- certainly there is a localized feature relatively close to the Sun, but the difficulties of observing through the confusion of the disc plane make it difficult to be absolutely certain whether it is global in extent. Further reversals have been claimed to exist, but with increasing uncertainty  \\citep[e.g.][]{sn79,b96,f01,menetal08,notkat10,kron11, vaneck11, beck11}. Speaking generally, the idea that magnetic field reversals are present in the Milky Way is widely accepted.  The situation in external galaxies \\cite[see the review][]{b96}, where the problem of observing from a position in the disc plane is absent, is nevertheless rather more uncertain. Except for the immediate neighbours of the Milky Way, the available resolution is too low to make definitive statements. But in the well observed nearby M31, reversals appear to be absent.  Classical mean-field dynamo models that start from a dynamically small field do not in general produce field reversals \\citep[e.g.][]{r88,m90}, because the leading eigenfunction of the kinematic mean-field galactic dynamo usually has no reversals, and this dominates the field structure as the field strength grows.  The situation becomes less straightforward when the seed field for the dynamo is stronger in comparison to the contemporary galactic magnetic field, so the galactic dynamo becomes nonlinear before the magnetic field configuration approaches the leading eigenfunction of the kinematic dynamo. It appears that long-lived reversals can survive in the form of nonlinear fronts \\citep{poezd,becketal94,metal98}: these are essentially field discontinuities, \\citep[see e.g.][]{betal94,vetal94,mps00,psm01}. However these models depend on rather carefully chosen initial conditions, and existing examples are one dimensional.  More recently, \\cite{metal12} have published an extended mean-field model in which long-lived reversals appear to be a quite common feature. This dynamo model differs in several important ways from earlier 2D mean-field models. Notably these computations begin from initial conditions of small-scale, approximately equipartition strength, fields that are randomly distributed in many discrete spots. This is in contrast to nearly all previous mean-field models, which begin from dynamically weak seed fields, either quasi-homogeneous or small-scale. The rationale is that these dynamically strong fields are the result of small-scale dynamo action in star forming regions in the protogalaxy. In the computations described in \\cite{metal12} there are also ongoing injections of small-scale field, but we will argue here that these are not central to the generation of reversals.  Our approach is close to that of the one dimensional model of \\cite{becketal94}, where a localized random initial condition was considered and long-lived large-scale magnetic reversals developed. (This contrasts with the model of \\cite{poezd}, where there were ongoing injections of random fields.) Thus both \\cite{poezd} and \\cite{becketal94} considered magnetic fields depending on galactocentric radius only, i.e. ring-like magnetic structures were described. The novelty of our approach is that we consider seed fields that depend on both radius and azimuth, i.e. there is no {\\it a priori} restriction to ring-like structures, and any such structures develop as a result of the dynamo process. We demonstrate below that ring-like structures with reversals, similar to those described by \\cite{becketal94}, also develop in our simulations. Thus our model reinforces and extends the results of \\cite{becketal94}.  On the other hand, we note that the stable magnetic ring-like structures with reversals obtained by \\cite{poezd} appeared in a model with some fine tuning of the model parameters. Stability conditions for such structures obtained by \\cite{betal94} depend on rather subtle properties of the galactic disc; this gave a hint that the magnetic configuration of the Milky Way with field reversals is a rare exception to the typical situation without reversals. Possibly, there was a misinterpretation of the results, and maybe the authors of the papers cited above were not insistent enough.  Nevertheless, dynamo generated configurations with reversals remained rather neglected by the astronomical community. Moreover, the option of including continuous injection of small-scale magnetic field as considered by \\cite{poezd}, was not further developed until recently.  In contrast, the new generation of dynamo models presented in \\cite{metal12} produce magnetic structures with reversals even for the simplest and most primitive distribution of the dynamo governing parameters and they appear to occur more or less as commonly as structures without reversals. This is why we consider below examples of dynamo models with (and without) reversals. They are demonstrably oversimplified and do not include any fine tuning of the governing parameters of the dynamo.  The paper \\cite{metal12} focussed on applications of the model to future observations of magnetic fields of the earliest galaxies using the recently developed (e.g. LOFAR) and planned (SKA) radio telescopes. The aim of this paper is to describe the long-lived reversals obtained as a quite general phenomenon of galactic dynamos. We are motivated here by future observations of magnetic fields in very young galaxies, whose detailed structures are almost unknown at the moment, and so we study a very simple, generic model. This is in some contrast to the approch of \\cite{poezd} (and to some extent \\cite{becketal94}) who focussed their attention on the dynamo properties of particular nearby galaxies (M31 and the Milky Way) and exploited particular subtle properties of their rotation curves and other parameters. ",
        "conclusions": "\\label{disc}  We deduce from these experiments that long-lived reversals can be generated by a standard 2D (nonaxisymmetric) mean-field dynamo model, given suitable initial conditions and a suitable range of the dynamo governing parameters. The initial conditions play a key role -- this is probably why such solutions have not been widely noted previously. Different steady solutions can be found for the same dynamo parameters, depending on the initial conditions. Large-scale reversals can appear after about $1$ Gyr, and they stabilize after $2-3$ Gyr. The important elements appear to be an inhomogeneous distribution of the initial field and strong differential rotation (larger $R_\\omega$), the latter effect being enhanced by initial fields that are locally near equipartition values. Note that even with larger values of $R_\\omega$ ($R_\\omega=20$ say), some initial conditions still produce final configurations without reversals.  We note that some of these results were anticipated by \\cite{becketal94}. This paper studied a one-dimensional model and showed that, with a sufficiently strong small-scale seed field, reversals could persist for times of order the age of the Universe.  As a straightforward consequence of the above conclusion, a systematic observational search for magnetic field reversals in a representative sample of spiral galaxies may provide valuable information concerning seed magnetic fields, and so clarify initial stages of galactic magnetic field evolution. This appears to be a realistic possibility for the forthcoming generation of radio telescopes. For a contemporary comparison, it might be appropriate for orientation to look at figure 11 of \\cite{vaneck11}, showing their best model of the reversals in the Milky Way field.  If we compare figure~\\ref{fig1}a with subsequent plots, the ongoing magnetic field injections gives an important additional mechanism to generate magnetic reversals. In addition to the global scale magnetic field reversals discussed above, this model contains several local reversals which appear to be due to local field injections, and which exist for some time. Possibly, such local reversals can be compared with reversals suggested to exist in the Milky Way \\citep[see e.g.][]{han}  We appreciate that the models presented are very oversimplified, but our intention is to show that initial conditions can strongly influence the long term evolution of the large-scale magnetic field. With \"standard\" weak initial fields ($|{\\bm B}| << B_{\\rm eq}$), any reversals are transient features. It appears important that the nonlinearity operates {\\it before} the large-scale field forms. Of course, confirmation or revision of the conclusions of this paper in the framework of 3D mean-field models, or even by direct numerical simulations of the microscopic induction equation, performed with realistic hydrodynamical models of very early galaxies would be an important development of the theory of galactic dynamos. However this obviously is beyond the scope of this paper.  We are grateful to our colleagues, who work on magnetic field evolution in very young galaxies namely T.~Arshakian, R.~Beck, M.~Krause, R.~Stepanov, for fruitful discussions which stimulated the writing of this the paper. We also thank Anvar Shukurov and an anonymous referee for critical readings of the paper."
    },
    "1208/1208.1997_arXiv.txt": {
        "abstract": "We present template radial velocity curves of $ab$-type RR Lyrae stars constructed from high-precision measurements of ${\\rm H\\alpha}$, ${\\rm H\\beta}$, and ${\\rm H\\gamma}$ lines. Amplitude correlations between the Balmer line velocity curves, Johnson $V$-band, and SDSS $g$- and $r$-band light curves are also derived. Compared to previous methods, these templates and derived correlations reduce the uncertainty in measured systemic (center-of-mass) velocities of RR Lyrae stars by up to 15 {\\kms}, and will be of particular interest to wide-area spectroscopic surveys such as the Sloan Digital Sky Survey (SDSS) and LAMOST Experiment for Galactic Understanding and Exploration (LEGUE). ",
        "introduction": "}  RR Lyrae stars are old ($\\gtrsim9$ Gyr), low-mass ($\\sim0.7M_\\sun$), pulsating stars ($V$-band amplitudes of $A_V\\sim1$ mag and periods of $\\sim 0.6$ days) that reside in the instability strip of the horizontal branch \\citep{smi95}. They are valuable tracers of old stellar populations and have recently been used to map the Galactic halo structure (i.e., its stellar number density profile) and substructure (e.g., tidal streams and halo overdensities; \\citealt{viv01}, \\citealt{kel08}, \\citealt{mic08}, \\citealt{wat09}, \\citealt{ses10a}).  While the spatial distribution of RR Lyrae stars is a powerful indicator of candidate halo substructures (e.g., Fig.~11 of \\citealt{ses10a}), it does not always reveal the full nature of halo substructures. For example, the Pisces Overdensity, a spatial group of about a dozen RR Lyrae stars at $\\sim80$ kpc from the Sun (clump ``J'' in \\citealt{ses07}; \\citealt{wat09}), was initially suspected to be a part of a tidal stream \\citep{ses10a}. However, spectroscopic observations by \\citet{kol09} and \\citet{ses10b} have revealed that the overdensity actually consists of {\\em two} velocity groups; one moving towards and the other one moving away from the Galactic center. This result suggests that the Pisces Overdensity may be a pile-up of tidal debris near the orbital turn-around point of a now-disrupted progenitor (which was, most likely, a dwarf spheroidal galaxy; \\citealt{ses10b}).  The \\citet{kol09} and \\citet{ses10b} findings underline the importance of obtaining kinematic data on candidate halo substructures; the nature and interpretation of a halo substructure can quickly change once kinematic data become available. Yet, relative to photometric data, which is abundant in this era of wide-area and multi-epoch surveys, spectroscopic observations are more scarce and may be considered as more ``expensive'' in terms of telescope time. Obtaining the kinematic information for variable stars, such as RR Lyrae stars, is even more complicated as one needs to subtract the velocity due to pulsations to get the center-of-mass (hereafter, systemic) velocity.  There are two approaches to this problem of determining the systemic velocity of RR Lyrae stars. In the first approach, one tries to observe RR Lyrae stars at a particular point during a pulsation period when the observed radial velocity is presumed to be equal to the star's systemic velocity (usually at ${\\rm phase\\sim0.5\\pm0.1}$, see Section~\\ref{systemic_velocity}). Unfortunately, sometimes it is very difficult to schedule observations to target this point (e.g., when executing observations in queue or service mode). In addition, RR Lyrae stars are dimmer near phase 0.5 ($\\sim0.3-0.5$ mag compared to maximum light at phase of 0), and thus require longer exposures to reach the same signal-to-noise ratio. The danger of longer exposures is blurring of spectral lines due to pulsations, which in turn decreases the precision of radial velocity measurements.  In the second approach, one assumes a certain model (a template) that describes the radial velocity due to pulsations as a function of pulsation phase. The amplitude, $A_{rv}$, of the radial velocity curve template (hereafter, template) is then scaled using a correlation between the amplitudes of velocity curves and light curves. The velocity at any phase can then be easily read-out once the scaled template is matched to one or several radial velocity observations taken at known phases. An example of such a template and an equation relating velocity and light curve amplitudes is provided by \\citet{liu91} (see his Equation 1). The advantage of this approach is that the systemic velocity can be estimated using radial velocities obtained in phases different from phase of 0.5.  Unfortunately, Liu's template and his velocity vs.~light curve amplitude relation are not really suitable for estimating systemic velocities of distant (i.e., faint) RR Lyrae stars because they were derived from measurements of {\\em metallic} lines and should only be used with measurements of such lines. This is a problem since measuring radial velocities using metallic lines requires high-resolution spectroscopy ($R\\gtrsim10000$), and obtaining quality, high-resolution spectra of faint ($V>18$) stars with exposures that are a small fraction of the pulsation period ($\\lesssim10\\%$ to avoid spectral blurring) is virtually impossible with today's facilities.  Instead of using weak metallic lines, a more practical choice would be to use strong Balmer lines (${\\rm H\\alpha}$ to ${\\rm H\\delta}$), which can be observed with more widely available low-resolution ($R\\sim1500$) spectrographs. However, Liu's template and his velocity vs.~light curve amplitude relation cannot be used with measurements of Balmer lines because the Balmer line velocity curves have a different shape {\\em and} amplitude than the velocity curve of metallic lines \\citep{oke62}. A method adopted by many studies (e.g., \\citealt{hb85, lay94, vzg05, pri09, ses10b, ses12}) is to use the radial velocity curve of RR Lyrae star X Ari measured from the ${\\rm H\\gamma}$ line by \\citet{oke66}. This is not an ideal solution because the velocity curve amplitude of X Ari ($A_{rv}\\sim85$ {\\kms}, see Figure~1 of \\citealt{oke66}) will likely not be similar to the velocity curve amplitude of some other star with a different $V$-band light curve amplitude (${\\rm A_V}$). The consequence of using an inadequate model for radial velocities is a more uncertain estimate of the systemic velocity.  The ability to precisely measure systemic velocities of RR Lyrae stars is very important as it may help with various Galactic studies, such as identifying halo substructures using kinematics (e.g., the Pisces Overdensity groups), discerning the nature of progenitors of halo substructures (e.g., globular cluster vs.~a dwarf spheroidal galaxy), and constraining orbits of halo substructures (e.g., the Virgo Overdensity; \\citealt{cd09, car12}). In addition, obtaining precise systemic velocities of RR Lyrae stars using one or two measurements would be beneficial to spectroscopic surveys such as the LAMOST Experiment for Galactic Understanding and Exploration (LEGUE; \\citealt{den12}), which may observe several thousands of RR Lyrae stars within 30 kpc of the Sun, and which plan to study the kinematics and structure of the Galactic halo.  In this paper, we use high-precision measurements of ${\\rm H\\alpha}$, ${\\rm H\\beta}$, and ${\\rm H\\gamma}$ lines of six field $ab$-type RR Lyrae stars to first construct template radial velocity curves (Section~\\ref{templates}). We then derive correlations between amplitudes of Balmer line velocity curves and light curves observed in Johnson $V$-band (Section~\\ref{correlations}) or SDSS $g$- and $r$-bands (Section~\\ref{sdss_correlations}). The systemic velocity of RR Lyrae stars and its uncertainty are defined in Section~\\ref{systemic_velocity}, and the results, along with implications for surveys such as LINEAR, PTF, LEGUE, and LSST, are discussed in Section~\\ref{discussion}. ",
        "conclusions": "}  We have presented template radial velocity curves of $ab$-type RR Lyrae stars constructed from high-precision measurements of ${\\rm H\\alpha}$, ${\\rm H\\beta}$, and ${\\rm H\\gamma}$ lines. Compared to the template radial velocity curve (solid line in the top right plot of Figure~\\ref{fig1}), the shape of observed ${\\rm H\\alpha}$ velocity curves varies the most (up to $\\sim10$ {\\kms}). The observed metallic lines curves show the least amount of variability ($\\sim3\\%$ of ${\\rm A_{rv}}$ or $\\sim2$ {\\kms} for a RR Lyrae star with ${\\rm A_V} = 1.0$ mag), followed by ${\\rm H\\gamma}$ and ${\\rm H\\beta}$ velocity curves ($\\sim3\\%$ to $4\\%$ of ${\\rm A_{rv}}$ or $\\sim4$ {\\kms}). The fluctuation in the shape of metallic, ${\\rm H\\gamma}$ and ${\\rm H\\beta}$ velocity curves is consistent with measurement errors, and implies little or no variation in the shape of velocity curves over a wide range of $V$-band amplitudes (0.6 to 1.1 mag), even for RR Lyrae stars that are undergoing Blazkho modulations.  We have found tight correlations (rms scatter $<3.4$ {\\kms}) between amplitudes of Balmer lines velocity curves (${\\rm A_{rv}}$) and $V$-band light curves (${\\rm A_V}$). A similar correlation, but for metallic lines, was first reported by L91. We also find that the ratio of amplitudes of Balmer versus metallic lines velocity curves decreases towards shorter wavelengths, and is 1.90, 1.54, and 1.39 for ${\\rm H\\alpha}$, ${\\rm H\\beta}$, and ${\\rm H\\gamma}$ lines. This pattern basically follows the depth of formation of lines; the ${\\rm H\\gamma}$ line forms closer to the photosphere than the ${\\rm H\\alpha}$ line, and therefore has smaller variations in the radial velocity.  The correlations derived in this paper have the potential to significantly improve the precision of systemic velocities of RR Lyrae stars estimated from a small number of radial velocity measurements of Balmer lines. For example, Equation~\\ref{Hgamma} indicates that previous studies that used the ${\\rm H\\gamma}$ velocity curve of X Ari (${\\rm A_{rv}\\sim85}$ {\\kms}, ${\\rm A_V\\sim1.0}$ mag) may have introduced up to $\\sim15$ {\\kms} of uncertainty in their estimates of systemic velocities for RR Lyrae stars with $V$-band amplitudes of 0.6 or 1.3 mag. Such uncertainties can be now be eliminated by using Balmer line radial velocity templates and Equations~\\ref{Halpha} to~\\ref{Hgamma} presented in this work.  We find the dominant source of uncertainty in the systemic velocity to be the phase at which the radial velocity is equal to the systemic velocity. We have estimated this uncertainty in phase to be about 0.1. For a RR Lyrae star with a $V$-band amplitude of ${\\rm A_V=1.0}$ mag, this uncertainty in phase translates into a velocity uncertainty of about 13 {\\kms}. A repeat of the analysis by \\citet{oke62} (see his his Section 6) on radial velocity data used in our work may provide more insight into this fundamental uncertainty.  Regrettably, we could not obtain radial velocity measurements of the ${\\rm H\\delta}$ line (due to low signal-to-noise ratio in echelle data; George W. Preston, private communication), and were thus unable to construct a template curve or establish a correlation between the ${\\rm H\\delta}$ velocity amplitude and the $V$-band light curve amplitude. This is quite unfortunate as this line is often used in radial velocity measurements of RR Lyrae stars and is usually accessible (along with ${\\rm H\\beta}$ and ${\\rm H\\gamma}$ lines) to most low-resolution spectrographs, such as DBSP \\citep{og82} or LRIS \\citep{oke95}. A spectroscopic survey that would provide velocity measurements of the ${\\rm H\\delta}$ line for several RR Lyrae stars would be very useful, and would allow the extension of this work to the ${\\rm H\\delta}$ line.  By deriving correlations between light curve amplitudes of RR Lyrae stars observed in SDSS $g$ and $r$ bands and the Johnson $V$ band (Section~\\ref{sdss_correlations}), we have expanded the applicability of Equations~\\ref{Halpha} to~\\ref{Hgamma} to a large number of already or soon to be observed RR Lyrae stars. For example, there are $\\sim380$ $ab$-type RR Lyrae stars with measured light curve parameters (period, epoch of maximum light, amplitudes in SDSS $ugriz$ band, etc.)~that have been observed in SDSS stripe 82 \\citep{ses10a}, a large fraction of which have spectroscopic observations \\citep{del08}. While the systemic velocities of these stars have already been measured \\citep{del08}, these measurements could be improved using Equations~\\ref{Halpha} to~\\ref{Hgamma}, leading to better estimates of kinematic properties of the Galactic halo.  Then there are multi-epoch, wide-area surveys such as the Lincoln Near-Earth Asteroid Research (LINEAR; \\citealt{ses11}) and Palomar Transient Factory (PTF; \\citealt{law09}), which are currently finding thousands of RR Lyrae stars within 30 kpc from the Sun \\citep{sesar11}. Some of these stars have previously been spectroscopically observed by SDSS, and many more are expected to be observed by the upcoming low-resolution spectroscopic survey LEGUE ($R=1800$; \\citealt{den12}). By using light curve parameters of RR Lyrae stars from LINEAR and PTF, radial velocities from LEGUE, and Equations~\\ref{Halpha} to~\\ref{Hgamma}, one could easily construct a sample of RR Lyrae stars with precise positions and 3D kinematics, and use them to search for halo substructures within 30 kpc from the Sun.  And finally, the Large Synoptic Survey Telescope (LSST; \\citealt{ive08}) is expected to be able to observe RR Lyrae stars as far as 360 kpc from the Sun \\citep{olu12}. With RR Lyrae stars detected at these distances, LSST will be able to search for halo streams and dwarf satellite galaxies within a significant fraction of the Local Group. Due to faintness of these distant RR Lyrae stars ($V\\lesssim23.4$), the spectroscopic followup will most likely involve observations of Balmer lines at low resolutions ($R\\sim1500$), making the tools presented in this work a natural choice for the measurement of their systemic velocities."
    },
    "1208/1208.6579_arXiv.txt": {
        "abstract": "We confront two types of phantom dark energy potential with observational data. The models we consider are the power-law potential, $V\\propto\\phi^{\\mu}$, and the exponential potential, $V\\propto \\exp\\left(\\lambda\\phi/M_P\\right)$. We fit the models to the latest observations from SN-Ia, CMB and BAO, and obtain tight constraints on parameter spaces. Furthermore, we apply the goodness-of-fit and the information criteria to compare the fitting results from phantom models with that from the cosmological constant and the quintessence models presented in our previous work. The results show that the cosmological constant is statistically most preferred, while the phantom dark energy fits slightly better than the quintessence does. ",
        "introduction": "}  Observations over the past dozen years have shown that the universe is currently under accelerating expansion (see~\\cite{Frieman:2008sn},~\\cite{Caldwell:2009ix} for reviews). Under the framework of general relativity and standard cosmology, a new form of exotic energy with negative pressure ($p<-\\rho/3$) is required to explain this phenomenon. Current observations suggested this so called \"dark energy\" made up about 73\\% of the energy density of the universe~\\cite{Suzuki:2011hu}~\\cite{Komatsu:2010fb}~\\cite{Blake:2011en}. So far $w_{DE}$ has been constrained to be very close to $-1$ assuming it is constant and the universe is flat. This seems to suggest the observation data prefer a cosmological constant as dark energy. However, a parametrized dark energy $w=w_0+w_a(1-a)$ has $\\sim100\\%$ uncertainty in dynamical parameter $w_a$ , when compared to data. This means dynamical dark energy models are not excluded.  Several dark energy model have been proposed in order to explain the cosmic acceleration. In addition to the most discussed cosmological constant, dynamically evolving scalar-field dark energy has been widely studied (for examples, see~\\cite{Ratra:1987rm}~\\cite{Caldwell:1998ii}~\\cite{Zlatev:1999tr}). Quintessence is a specific case of scalar-field dark energy with canonical kinetic terms, which admits $-1<w_{\\phi}<1$. It has drawn much attention because it can in principle provide the \"tracker\" property -- a property that makes the energy density today insensitive to its initial condition, i.e., the initial energy density value $\\rho_{\\phi}$ ~\\cite{Zlatev:1999tr},~\\cite{Steinhardt:1999nw}. While fine-tuning of potential parameter is still required~\\cite{Singh:2003vx}, tracker quintessence can alleviate the cosmic coincidence problem because a precise setting of initial energy density ratio for matter and dark energy is no longer required.  Present observational constraint on $w_{DE}$ still allows for $w_{DE}<-1$. For example, results from WiggleZ ~\\cite{Blake:2011en} showed current constraints on constant $w$ is $-1.114<w<-0.954$ after combining with the latest SN-Ia, CMB and BAO data. We note that while quintessence only allows for$-1<w_{\\phi}<1$, phantom scalar-field, another dynamical DE model, proposed by Caldwell, Carroll et al~\\cite{Caldwell:1999ew}, that invokes a negative kinetic energy, can satisfy $w_{\\phi}<-1$. While there exist several known theoretical difficulties for phantom scalar-field dark energy model such as the violation of null dominant energy condition (NDEC)~\\cite{Carroll:2003st} and a possible Big Rip phase in the future~\\cite{Caldwell:2003vq}~\\cite{Scherrer:2004eq}, it is still very worth while to confront it directly and independently against the cosmological observations, especially since the current best-fit for equation-of-state is smaller than $-1$(see~\\cite{Caldwell:2003vq}~\\cite{Guo:2004ae}~\\cite{Sun:2004bf} for examples).  In our previous work~\\cite{Wang:2011bi}, we have tested several tracker quintessence models with observational data. The result showed that the best-fit of the inverse-power-law potential and the inverse-exponential potential models both reduced to the cosmological constant. Motivated by this and the implication of $w_{DE}<-1$ from observations, in this paper we put the phantom dark energy models to the test. We consider two specific scalar-field phantom potentials: the power-law potential~\\cite{Guo:2004ae} and the exponential potential~\\cite{Li:2003ft}, each with one free parameter. The reason to choose these potentials is that they also possess the attractor-like property: insensitive to initial conditions. We take the model-based approach to confront the models with observational data. The data we use includes the latest Type-Ia supernova (SN-Ia) compilation set, the cosmic microwave background (CMB) and the baryon acoustic oscillations (BAO) observations. We also confront these models with the cosmological constant and the quintessence scalar-field models by using the Goodness-of-Fit test and the information criteria to assess the merit of each model. ",
        "conclusions": "We have tested two potential forms of phantom scalar-field dark energy, the power-law $V=V_1\\left(\\phi/M_P\\right)^{\\mu}$ and the exponential potential, $V=V_2 \\exp\\left(\\lambda\\phi/M_P\\right)$, with current observations. Tight model parameter constraints are obtained in Table~\\ref{table:results}. We also provided results of cosmological constant and two types of quintessence potentials, the inverse-power law potential and the inverse-exponential potential. Cosmological constant yields the best goodness-of-fit and smallest information criteria. This shows that the cosmological constant is still the most preferred dark energy model among all that we have considered, even with the constraint $w>-1$ removed. Phantom models fit worse to observations than that with the cosmological constant, but are slightly better than the two quintessence potential models. Although the current observations still prefer the cosmological constant as the dark energy, phantom and quintessence models under considerations are only slightly worse in terms of GoF and the information criteria; all our models in Table~\\ref{table:results} have GoF$\\sim85\\%$, $\\Delta$BIC$\\sim6$, and $\\Delta$AIC$\\sim2$.  Another interesting result is the best-fit for phantoms. Unlike quintessence, which have the best-fit equivalent to $\\Lambda$CDM~\\cite{Wang:2011bi}, that for phantom models moves away from $w_{\\phi}=-1$ (see Fig.~\\ref{wbest}), as the best-fit for parameters $\\lambda$ and $\\mu$ are no longer zero (Fig.~\\ref{joint}). For $V=V_1\\left(\\phi/M_P\\right)^{\\mu}$, the best-fit $w_{\\phi}(0)\\sim -1.03$; as for $V=V_2 \\exp\\left(\\lambda\\phi/M_P\\right)$, we find $w_{\\phi}(0)\\sim -1.06$ for the best fit. This indicates that the current observations may prefer $w_{DE}<-1$.  In summary, the model-based approach we used in this paper suggests that the cosmological constant is more preferred, and dark energy models with $w<-1$ is preferred over $w>-1$. Notice that in both phantom models the cosmological constant cases are still inside $1-\\sigma$ range. This means currently we cannot distinguish small-dynamical dark energy from the cosmological constant. Future observations from next generation dark energy probes are expected to constrain $w$ about ten times better than the present value~\\cite{astro-ph/0609591}. More stringent constraints on the parameter space are thus expected to be obtained. By then we should be able to attain more insights into the physics of dark energy models with this model-based approach, or even rule out some of the models at a sufficient confidence level (see results with projected data in~\\cite{Yashar:2008ju}~\\cite{Barnard:2007ta}~\\cite{Bozek:2007ti}~\\cite{Abrahamse:2007te})."
    },
    "1208/1208.2099_arXiv.txt": {
        "abstract": "A moment approach for orbit determinations of an astrometric binary with low signal-to-noise ratio from astrometric observations alone is proposed, especially aiming at a close binary system with a short orbital period such as Cyg-X1 and also at a star wobbled by planets. As an exact solution to the nonlinearly coupled equation system, the orbital elements are written in terms of the second and third moments of projected positions that are measured by astrometry. This may give a possible estimation of the true orbit. ",
        "introduction": "Space astrometry missions such as Gaia and JASMINE are expected to reach a few micro arcseconds \\citep{GAIA-1,GAIA-2,JASMINE}. Moreover, high-accuracy VLBI is also available.  For visual binaries, formulations for orbit determinations have been well developed since the nineteenth century \\citep{Thiele,Binnendijk,Aitken,Danby,Roy}. At present, numerical methods are successfully used \\citep{EX,CO,OC}. Furthermore, an analytic solution for an astrometric binary, where one object can be observed and the other such as black holes and neutron stars is unseen, has been found \\citep{AAK,Asada2008}. The solution requires that sufficiently accurate measurements of the position of a star (or the photo-center of a binary) are done more than four times for one orbital period of the binary system.  For a close binary system with a short orbital period, we have a relatively large uncertainty in the position determination. For instance, the orbital period of Cyg-X1 is nearly 6 days, which are extremely shorter than that of normal binary stars, say a few months and several years. Because of such an extreme condition, it is interesting to seek another method in addition to the standard one. Moreover, stars with planets also are another interesting target.  What can we do for orbit determination from position measurements with low signal-to-noise (SN) ratio? It is expected that the position of the object is measured many times. The dense region of the observed points is corresponding to the neighborhood of the apastron of the Kepler orbit, because the motion of the source star is slower according to the Kepler's second law. On the other hand, a region of fewer points is including the periastron, around which the source star moves faster. Therefore, a statistical analysis including the variance of the measured positions and their correlation will bring the information about the orbital elements of the binary system. Gedanken experiments suggest that the second moments are useful for exploring the shape of the orbit but they are not sufficient for the full orbit determination. At least the third moments seem to be needed. See Figs. \\ref{moment-1} and \\ref{moment-2}.  Therefore, the main purpose of this paper is to propose a method for orbit determination of an astrometric binary with low SN ratio by using the second and third moments. We shall provide also an exact solution for the coupled equations. As a result, the orbital elements of the binary are written in terms of the second and third moments.      This paper is organized as follows. We present a formulation and the solution in $\\S$ 2. In $\\S$ 3, numerical tests are also done to see how reliable the analytic result is for practical cases taking account of observational noises. $\\S$ 4 is devoted to Conclusion. Throughout this paper, the spatial coordinates are the angular positions normalized by the distance to the celestial object. ",
        "conclusions": "This paper proposed a moment approach for orbit determinations of a close binary system with a short orbital period from astrometric observations alone. As an exact solution to the coupled equations, the orbital elements are written in terms of the second and third moments of the projected position that is measured by astrometry.  The moment formalism does not replace the standard method using Kepler equation. It is safer to say that the present formalism is a supplementary tool for giving a rough parameter estimation, which can be used as a trial value for full numerical data fittings. It is interesting to make numerical tests of the present method. It is left as a future work.  In the moment approach, the temporal information is smeared. Therefore, the orbital period cannot be determined by this method. Another method such as the Fourier analysis of position data with time (in Fig. \\ref{points} for example) would give a characteristic frequency that is the inverse of the orbital period. Fourier analyses recover the orbital period from numerically simulated data for Fig. \\ref{points}. This suggests that the moment approach can be applied also to unknown binary systems, if a Fourier analysis is adequately used to know the orbital period. Namely, the method could be used to search new binary systems.   We would like to thank Professor N. Gouda for useful information on astrometry missions. We wish to thank the JASMINE science WG member for stimulating conversations. We would be grateful to Y. Sendouda, R. Takahashi and K. Izumi for helpful conversations on numerical simulations. This work was supported in part (H.A.) by a Japanese Grant-in-Aid for Scientific Research from the Ministry of Education, No. 21540252 (Kiban-C) and in part (K.Y.) by JSPS research fellowship for young scientists."
    },
    "1208/1208.0780_arXiv.txt": {
        "abstract": "IRAS 20050+2720 is young star forming region at a distance of 700~pc without apparent high mass stars. We present results of our multiwavelength study of IRAS 20050+2720 which includes observations by \\emph{Chandra} and \\emph{Spitzer}, and 2MASS and UBVRI photometry. In total, about 300 YSOs in different evolutionary stages are found. We characterize the distribution of young stellar objects (YSOs) in this region using a minimum spanning tree (MST) analysis. We newly identify a second cluster core, which consists mostly of class~II objects, about 10\\arcmin{} from the center of the cloud. YSOs of earlier evolutionary stages are more clustered than more evolved objects.  The X-ray luminosity function (XLF) of IRAS 20050+2720 is roughly lognormal, but steeper than the XLF of the more massive Orion nebula complex. IRAS 20050+2720 shows a lower $N_H/A_K$ ratio compared with the diffuse ISM. ",
        "introduction": "Star formation occurs in dense and cool molecular clouds, which collapse under their own gravity. We observe those clouds over a range of masses, but the majority of stars forms in clusters with more than a hundred members \\citep{2000AJ....120.3139C,2003ARA&A..41...57L,2003AJ....126.1916P}. Thus, studying star formation in those clusters is an important step to understand the history of the stars we currently observe, and the stellar population of the outer galaxy as a whole \\citep[for general reviews of star formation see][]{2007prpl.conf.....R}.  In the early stages the proto-star is hidden by its parent circumstellar material and thus can only be observed as far-IR emission from the warm, dusty envelope (class 0). This envelope becomes less dense in class I objects. They often drive powerful outflows and carve holes in their envelopes so radiation can escape. Eventually, the envelope disperses and the star becomes visible as a classical T~Tauri star (CTTS) or ``class II'' source in IR classification \\citep{1987IAUS..115....1L}. The disk still causes an IR excess over photospheric stellar emission. On the star itself accretion shocks and coronal activity lead to X-ray emission \\citep{lamzin,2002ApJ...567..434K,acc_model,2011AN....332..448G}. Later, the IR excess vanishes as the gas disk is dissipated \\citep{2012arXiv1205.2049I}. In this stage (weak-line T~Tauri stars = WTTS) cluster members cannot be distinguished from main-sequence (MS) stars by IR observations. One method to identify them is through their X-ray luminosity, which is far higher than for most older stars \\citep[for a review of X-ray properties of young stellar objects (YSOs) before high-resolution X-ray spectroscopy with \\emph{Chandra} or \\emph{XMM-Newton} see][]{1999ARA&A..37..363F}.   Many details of the star formation process in clusters are still under debate despite significant observational progress. For example, it is unclear if in an undisturbed cluster high-mass or low-mass stars form first. However, the theoretical assumption of a cluster evolving with minimal influence from the rest of the galaxy may be far from realistic. For example, in $\\rho$~Oph \\citet{1992A&A...262..258D} suggested that the expanding shell of the upper Sco star forming region travelled through the cloud and triggered the star formation. Events like that would cause a non-uniform age distribution within a cluster.  In this paper we study the young cluster IRAS~20050+2720, which is located at a distance of 700~pc based on CO velocities \\citep{1989ApJ...345..257W}. It is found at a galactic longitude where the Cygnus Rift and the Cygnus X region are distinguishable, therefore, there is no distance ambiguity.  We use this distance throughout the article. There appear to be no massive stars in IRAS 20050+2720, thus the intensity of the ambient radiation should be small and we can study the evolution of late-type stars in the absence of external irradiation on their disks.  IRAS~20050+2720 is associated with a water maser \\citep{1991A&A...246..249P} and a methanol maser \\citep{2010A&A...517A..56F}. Molecular line emission indicates mass infall \\citep{1997ApJ...484..256G,1999ApJS..122..519C}. Radio observations reveal a complex structure of several lobes close to the nominal source position of IRAS~20050+2720, which has been identified as two or three jets from different proto-stellar sources \\citep{1995ApJ...445L..51B,1999A&A...343..585C,2008A&A...481...93B}. % The luminosity of this region is estimated as 388~$L_{\\sun}$ from the \\emph{IRAS} fluxes \\citep{1996A&A...308..573M}. \\citet{2001A&A...369..155C} obtained mm-maps of the regions, estimatinge a total gas mass of 200~$M_{\\sun}$ within 65\\arcsec{} (0.2~pc) radius.     The first attempt to establish an IR classification of individual stars in IRAS~20050+2720 identified about 100 cluster members, about half of which show an IR excess \\citep{1997ApJ...475..163C}. While they estimate an average cluster age of 1~Myr, they interpret the radio lobes and the strong reddening as signatures for a more recent star formation event in the last 0.1~Myr. \\citet{2005ApJ...632..397G} confirmed these ideas. %  With \\emph{Spitzer} it is possible to get very accurate photometry of star forming regions and classify individual objects according to their IR spectral energy distribution (SED). \\citet{2009ApJS..184...18G} identified 177 YSOs in IRAS~20050+2720, which they grouped in two distinct cores. Compared with \\citet{1997ApJ...475..163C} they identify more absorbed cluster members and show that the earlier definition of subclusters is partially due to an observational bias. In this paper, we extend the spatial coverage with an additional \\emph{Spitzer} observation and add information from a \\emph{Chandra} observation and, surprisingly, the first dedicated optical photometry.  Our goal is to compare the clustering and the X-ray properties of IRAS 20050+2720 with other star forming regions and to search for internal gradients of density and age that could indicate triggered star formation.  In Sect.~\\ref{sect:obs} we describe the observations of IRAS~20050+2720. In Sect.~\\ref{sect:analysis} we identify disk-bearing YSOs by the IR colors and disk-less YSOs based on their X-ray emission. Our results in Sect.~\\ref{sect:results} describe the spatial distribution of YSOs and the optical photometry allows us to place individual members on evolutionary tracks. We show the X-ray properties and discuss the $A_K/N_H$ ratio. We end with a short summary in Sect.~\\ref{sect:summary}. ",
        "conclusions": "\\label{sect:summary} We present multi-wavelength data for the young star forming region IRAS~20050+2720 from the \\emph{Spitzer} and \\emph{Chandra} satellites and the ground-based 2MASS and IPHAS surveys as well as targeted FLWO photometry. YSOs of class~I and II are identified based on their IR properties. X-ray sources with optical or IR counterparts that are compatible with the cluster distance are added to this sample as well as a group of X-ray sources that are identified as cluster members based on their close association with the densest cluster core.  We use a minimum spanning tree (MST) analysis to characterize the spatial distribution of sources in IRAS 20050+2720. Overall, class~I sources are more clustered than class~II sources, and class~II sources are more clustered than class~III sources. %  Two main cluster cores can be identified in IRAS 20050+2720. Cluster core~E consists only of class~II and III sources, but seems younger in the $r'$ vs. $r'-i'$ diagram. It is surrounded by an IR nebulosity.  Cluster core~W has a dense and highly extincted core in the South and a much less dense region in the north. It is not clear if those two regions are distinct groupings along overlapping lines-of-sight, or if they are part of the same physical cluster core. The age of IRAS 20050+2720 is $<3$~Myrs and it might contain populations of different ages.  In X-rays we detect preferentially the more massive YSOs; the detection fraction is higher for more evolved evolutionary stages, where the absorbing column densities are lower. Still, the detected class~I objects are both hotter and more luminous than class~II sources with similar absorption. We compare the XLF with other clusters and find an apparent correlation with cluster mass: The XLF can always be approximated by a lognormal distribution, but it seems to be steeper for clusters of lower mass.  Using the combined X-ray and infrared data, we derive the $N_H/A_K$ values for the YSOs in IRAS 20050+2720.  We find a value which is 60\\% of the value for the diffuse ISM with no significant difference in this ratio between the cluster cores of IRAS 20050+2720 or between sources at different evolutionary stages. Compared to values derived with other YSOs the $N_H/A_K$ ratio is higher than that for the low-mass regions Serpens or NGC~1333, but much lower than that observed for class~II objects in the Taurus molecular cloud where the extinction and X-ray absorption may be dominated by circumstellar material. This could be due to modest grain growth in the molecular cloud environment that dominates the extinction toward IRAS 20050+2720; in contrast, the circumstellar environments may have elevated values of $N_H/A_K$ due to grain growth, gravitational settling, planet formation or dust evaporation."
    },
    "1208/1208.0113_arXiv.txt": {
        "abstract": "We report new features of the typical mixed-morphology (MM) supernova remnant (SNR) W44. In the X-ray spectra obtained with Suzaku, radiative recombination continua (RRCs) of highly ionized atoms are detected for the first time. The spectra are well reproduced by a thermal plasma in a recombining phase. The best-fit parameters suggest that the electron temperature of the shock-heated matters cooled down rapidly from $\\sim1$\\,keV to $\\sim 0.5$\\,keV, possibly due to adiabatic expansion (rarefaction) occurred $\\sim20,000$ years ago. We also discover hard X-ray emission which shows an arc-like structure spatially-correlated with a radio continuum filament. The surface brightness distribution shows a clear anti-correlation with $^{12}$CO ($J=2$--1) emission from a molecular cloud observed with NANTEN2. While the hard X-ray is most likely due to a synchrotron enhancement in the vicinity of the cloud, no current model can quantitatively predict the observed flux. ",
        "introduction": "\\label{sec:intro} Mixed-morphology (MM) supernova remnants (SNRs) are classified based on the characteristic morphology of centrally-filled thermal X-ray emission with a synchrotron radio shell (\\cite{Rho98}; \\cite{Lazendic06}; \\cite{Vink12}). More than 25\\% of the X-ray-detected Galactic SNRs are categorized as the MM type \\citep{Jones98}. Most of the MM SNRs indicate evidence for interaction with dense interstellar medium (ISM), CO and H\\emissiontype{I} emissions or OH (1720\\,MHz) masers in some cases \\citep{Rho98}. They are also associated with TeV/GeV $\\gamma$-ray sources, suggesting the presence of molecular clouds in their vicinity (e.g., \\cite{Albert07}; \\cite{Aharonian08}; \\cite{Hewitt09}; \\cite{Abdo10}). Such ambient environment would play some roles in the evolutions of the morphology and plasma condition in these SNRs; e.g., the central X-ray emission might be enhanced by cool clouds left relatively intact after the passage of blast waves to slowly evaporate in the hot SNR interior \\citep{White91}, or density gradients in the pre-existing ISM could contribute the emission profile \\citep{Petruk01}.  Recent Suzaku observations revealed the presence of enhanced radiative recombination continua (RRCs) in X-ray spectra of several MM SNRs, IC\\,443 \\citep{Yamaguchi09}, W49B \\citep{Ozawa09}, G\\,359.1$-$0.5 \\citep{Ohnishi11} and W28 \\citep{Sawada12}. The strong RRC emission can be observable when an ionization temperature ($kT_{\\rm z}$\\footnote{$kT_{\\rm z}$ represents populations of ionization states. In collisional ionization equilibrium, $kT_{\\rm z}$ is equal to the plasma temperature; see also \\citet{Masai94}.}) is significantly higher than electron temperature ($kT_{\\rm e}$) so that the free-bound transition (recombination process) becomes dominant---hereafter, we call it a recombining plasma (RP). These discoveries are dramatically changing our understanding of the SNR evolution, since most young or middle-aged SNRs have either an ionizing plasma (IP: $kT_{\\rm e}>kT_{\\rm z}$) or a collisional ionization equilibrium plasma (CIE: $kT_{\\rm z}=kT_{\\rm e}$). While hints of RPs were reported from a few non-MM SNRs (\\cite{Hughes03}; \\cite{Broersen11}), all the  robust results have been found only in MM SNRs. It is suggested, therefore, that the formation process of the RPs is closely related to the characteristic environment of the MM SNRs.  The environment may also affect the particle acceleration in SNR blast waves. The shock waves propagating into the inhomogeneous ambient medium may produce strong turbulence in the magnetic field, which enhances the efficiency of the particle acceleration in this region (e.g., \\cite{Giacalone07}; \\cite{Inoue12}).  \\begin{table*}[!t] \\caption{Observation logs.}\\label{tab:obs} \\begin{center} \\begin{tabular}{lcccccc} \\hline \\hline & Name & Obs. ID  & Obs. Date & (R.A., Dec.) $_{J2000}$ & ($l$, $b$) &Exposure\\\\ \\hline Source & W44  & 505004010   & 2010-Apr-10 &  (\\timeform{18h56m08s}, \\timeform{+01D23'19\"})  & (\\timeform{34.7D}, \\timeform{-0.41D}) & 61.1\\,ks\\\\ Background & GR & 500009020   & 2007-Aug-18 &  (\\timeform{18h44m01s}, \\timeform{+04D04'41\"})  & (\\timeform{28.5D}, \\timeform{+0.21D}) & 98.9\\,ks\\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  In this paper, we focus on one of the most typical MM SNRs, W44 (also known as G\\,34.7$-$0.4 or 3C\\,392) to study its X-ray properties and their relation with the environment. W44 is a middle-aged ($\\sim20,000$\\,yr; \\cite{Wolszczan91}) SNR located on the Galactic plane at a distance of $D\\sim3$\\,kpc (Claussen et al.\\ 1997). In the radio band, W44 shows a distorted shell of the synchrotron emission (e.g., \\cite{Kundu72}; \\cite{Handa87}). Recent high-resolution radio observations revealed more complicated filaments in the radio continuum shell (\\cite{Giacani97}; \\cite{Castelletti07}). Molecular radio emission was detected from the vicinity of W44 by several observations of the CO (\\cite{Wootten77}; \\cite{DeNoyer79}; \\cite{Seta98}; \\cite{Reach05}) and H\\emissiontype{I} \\citep{Koo95} lines. More recently, \\citet{Seta04} showed a whole view of the high-resolution $^{12}$CO emission map with the 17$\\arcsec$-beam of the Nobeyama 45-m radio telescope and reported many hints for the shock-cloud interaction, including cloud evaporation. The presence of 1720\\,MHz OH maser emission is also a sign of the interaction between the SNR shocks and the ambient molecular clouds (e.g., \\cite{Claussen97}; \\cite{Frail98}).  In the X-ray band, W44 was found to have center-filled morphology by ROSAT observations \\citep{Rho94}. Based on the Chandra result, \\citet{Shelton04} proposed that the centrally bright X-ray morphology is caused by either entropy mixing in the form of thermal conductions, or evaporation of swept-up clouds. The analyses of combined X-ray spectra of ROSAT, EXOSAT (ME), Einstein (SSS), and Ginga revealed that the plasma was IP and has not reached CIE, and the spectra are interpreted  by cloud evaporation (\\cite{Rho94}; \\cite{Harrus97}). ASCA obtained higher-quality spectra and determined the ionization temperature ($kT_{\\rm z}$) of Si from the flux ratio of K$\\alpha$ lines between H-like and He-like ions (H-like/He-like). The result was that $kT_{\\rm z}$ is nearly equal to the electron temperature $kT_{\\rm e}$ \\citep{Kawasaki05}, or the thermal plasma was nearly in CIE.  This paper reports revised X-ray features of W44 made with the Suzaku observation. The bright X-ray flux and apparent X-ray size of $\\sim\\timeform{0.6D}$ in diameter are best suited to study the detailed plasma structure. The high sensitivity and the good energy resolution of Suzaku enable us to determine the spatial structure of the X-rays accurately, in particular, that of the thermal plasma. We report the first detection of an RP in this SNR and discovery of hard X-ray emissions. The evolution of the RP is quantitatively examined, and then some scenarios for its origin are addressed. The hard X-ray structure is compared to that of the molecular cloud with the NANTEN2 observations. The origin is discussed in the context of a shock-cloud interaction. Throughout this paper, the distance to the SNR is assumed to be 3\\,kpc, and the errors quoted are at the 90\\% confidence level.   \\begin{figure}[!t] \\begin{center} \\FigureFile(85mm,85mm){figure1a.eps} \\FigureFile(70mm,70mm){figure1b.eps} \\end{center} \\caption{\\textit{Top Left}: Three-color XIS image of W44 overlaid on the Galactic coordinate. The red, green and, blue correspond to the 1.0--2.0\\,keV, 2.0--5.0\\,keV and 5.0--7.0\\,keV bands, respectively. The white contours represent X-ray surface brightness obtained by the ASCA GIS observation \\citep{Kawasaki05}. The FoV of the XIS is shown with the white square. The white cross indicates the location of the pulsar PSR\\,B1853$+$01. \\textit{Top Right}: Hard X-ray band  (5.0--7.0\\,keV) image obtained with the Suzaku XIS. The white contours represent the 1442.5\\,MHz radio continuum emission with VLA. \\textit{Bottom}: Example of the source spectrum extracted from the region C. The solid line is the CXB model spectrum.}\\label{fig:image} \\end{figure} ",
        "conclusions": "\\subsection{Recombining Plasma} We found that the spectra from all the regions of W44 were nicely fitted with the model of the NEIJ and power-law components. In order to clarify the property of the plasma state, we estimate the average charge per element and compared to those expected in a CIE plasma with the same electron temperature. The results derived with the best-fit parameters for the region A spectrum are shown in table \\ref{tab:average}. We see that the averaged charge states of lighter elements, C, N, and O are consistent with those in CIE, while those of the heavier elements, such as  Si, S and Fe are higher than CIE, indicating that these elements are in the recombining state. In contrast, some previous X-ray studies claimed that the plasma in W44 is ionizing or nearly in CIE (\\cite{Rho94}; \\cite{Kawasaki05}). One of the possible origins of this discrepancy is a difference in the energy resolution and/or photon statistics in the hard X-ray band.  Our result suggests that either  only the initial electron temperature $kT_{\\rm e1}$ cooled down rapidly to $kT_{\\rm e2}$ (here, electron cooling), or only the initial $T_{\\rm z}$ (=$kT_{\\rm e1}$) was raised higher than $kT_{\\rm e2}$(here, selective ionization). These led the ions to be higher charge states (RP)  compared to those in CIE. In the next sections, we separately discuss these two cases.  \\subsection{Electron Cooling} One possible scenario for the electron cooling is thermal conduction, so we first evaluate this possibility. In this scenario, the energy transfer from the central region ($kT_e = 0.48 \\pm 0.01$\\,keV) to the cooler periphery ($kT_e = 0.40^{+0.02}_{-0.01}$\\,keV) is considered. A scale length of the temperature gradient $l_{\\rm T}$ = $({\\rm grad} \\, ln T)^{-1}$ is calculated to be $1.5\\times10^{20}$\\,cm, at the distance to the SNR of 3\\,kpc. Therefore, the thermal conduction timescale $t_{\\rm{cond}}$ is estimated to be \\begin{eqnarray} t_{\\rm{cond}} \\simeq 1\\times10^7 \\left(\\frac{n_{\\rm e}}{\\rm1.0\\,cm^{-3}}\\right)\\left(\\frac{kT_{\\rm e{\\rm{(in)}}}}{0.48\\, \\rm{keV}}\\right)^{-5/2} \\nonumber \\\\ \\times \\left(\\frac{l_{\\rm T}}{\\rm 1.5\\times10^{20} \\, cm}\\right)^{2} \\rm{yr}. \\label{eq:t_cond} \\end{eqnarray} (\\cite{Spitzer62}; \\cite{Kawasaki02}), where the electron number density $n_{\\rm e}$ is calculated from the best-fit EM we obtained. The derived timescale is nearly 50 times longer than the age of W44 ($\\sim 20,000$\\,yr),  and hence the thermal conduction scenario is difficult to reproduce the observed electron temperature gradient.  The other possibility of the electron cooling is adiabatic expansion, which may occur in a rarefaction process; a rapid cooling of electrons occurs when a blast wave in a dense circumstellar medium (CSM) breaks into a lower-density ISM (\\cite{Itoh89}; \\cite{Shimizu12}; \\cite{Moriya12}). \\citet{Sawada12} argued the recombination time $t_{\\rm rec}$ is an essential parameter to evaluate the possibility of this scenario. We hence estimate the $t_{\\rm rec}$ using the best-fit $n_{\\rm e}t$ values: the lowest  at the east region of $3.8\\pm0.1\\times10^{11}$\\,s\\,cm$^{-3}$, and the average in the center region of $\\sim7.1\\pm0.1\\times10^{11}$\\,s\\,cm$^{-3}$ (figure \\ref{fig:map}-c). Since the electron number densities $n_{\\rm e}$ at the center and east region are, respectively $\\sim1.0$\\,cm$^{-3}$ and $\\sim0.8$\\,cm$^{-3}$, the recombining time is derived to be $\\sim$15,000--20,000\\,yr. \\citet{Wolszczan91} estimated the age of the pulsar PSR\\,B1853$+$01, physically associated with W44, to be 20,000\\,yr, which is consistent with our estimated recombination time. Therefore, the adiabatic cooling scenario is preferable in our case, because the shock-heating and the rarefaction should occur at the initial phase of the SNR evolution. \\citet{Broersen11} derived conditions for general SNRs to keep RP in their evolutions. Comparing the $t_{\\rm rec}$ with an adiabatic cooling time $t_{\\rm ad}$ in different $n_{\\rm e}$, they concluded that a low $n_{\\rm e}$ significantly decreases the recombination rate. They claimed SNR\\,0506$-$68 ($\\sim4,000$\\,yr) to be still RP in a dense ISM ($n_{\\rm e}\\sim10$\\,cm$^{-3}$). Applying their argument to W44's case, we confirmed that the $t_{\\rm rec}$ for W44 is about one order of magnitude above that for SNR\\,0506$-$68 (see also \\cite{Vink12}). This estimation also supports the adiabatic cooling scenario.   \\begin{table}[t] \\caption{Average Charge Per Element}\\label{tab:average} \\begin{center} \\begin{tabular}{ccc} \\hline \\hline Element & \\multicolumn{2}{c}{Average Charge} \\\\ & NEIJ (Region A) & CIE at 0.45\\,keV \\\\ \\hline C      \t&    \t\t5.99 \t&      \t5.99 \\\\ N      \t&          \t6.98 \t&      \t6.98 \\\\ O       \t&          \t7.92 \t&      \t7.92 \\\\ Ne     \t&           \t9.51 \t&        \t9.34\\\\ Mg    \t&          \t10.87 \t&        \t10.32\\\\ Si       \t&         \t12.26 \t&        \t12.01\\\\ S       \t&         \t13.95 \t&       \t13.91\\\\ Fe     \t&           \t16.45 \t&      \t16.37\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\subsection{Selective Ionization} Possible mechanisms to rise up $kT_{\\rm z}$ (selective ionization)  would be either ionization by supra-thermal electrons (e.g., \\cite{Tanaka86}; \\cite{Masai02}), or photo-ionization (e.g., \\cite{Piro00}). In W44, the 2--10\\,GeV and  0.4--3\\,GeV emissions were discovered with the Fermi LAT \\citep{Abdo10} and AGILE \\citep{Giuliani11}, respectively. The idea of the former scenario is that supra-thermal particles are injected into a Fermi acceleration process, and are accelerated up to very high energy. The supra-thermal particles also ionize atoms in a thermal plasma to higher ionized states, and produce the RP. In fact, large solar flares often show a signature of RP (in Fe atoms) and hard X-ray emissions \\citep{Tanaka86}. Thus, the scenario of supra-thermal particle ionization may predict positive correlation between the ionization temperature ($kT_{\\rm z}$) and the flux of TeV/GeV $\\gamma$-rays/hard X-rays. However, our observation shows no systematic spatial variation of $kT_{\\rm z}$, and hence possibility of selective ionization by supra-thermal particles is unlikely.  Another possibility, photo-ionization, requires a strong X-ray source in or near W44, but none of such sources is found at present. Therefore, such ionizing source should, if any, be bright in past. A $\\gamma$-ray burst is one of the candidates. If W44 was a remnant of hyper-nova explosion associated with a $\\gamma$-ray burst, the X-ray afterglow would be strong enough to ionize a substantial amount of atoms (high $kT_{\\rm z}=kT_{\\rm e1}$), whereas the electron temperature was relatively low (low $kT_{\\rm e2}$). A weak point of this scenario is, however, that no particular signature for a hyper-novae has so far been observed in W44.  \\subsection{Origin of Hard X-ray Emission}\\label{sec:hardxray} We detected strong hard X-rays from the west region of W44 (region C), and possibly from the entire SNR regions. A conceivable origin is a pulsar wind nebula (PWN) associated with PSR\\,B1853$+$01 (\\cite{Wolszczan91}; \\cite{Petre02}) located just outside the FoV (the white cross in figure \\ref{fig:image}). The Chandra ACIS found that the X-ray emitting region of the PWN is within $\\sim\\timeform{2'}$ from PSR\\,B1853$+$01, while we found that the peak of the hard X-rays is far from the PWN ($\\sim\\timeform{12'}$). The total flux of the PWN plus PSR\\,B1853$+$01 is $\\sim1.7\\times10^{-13}$\\,ergs\\,cm$^{-2}$\\,s$^{-1}$ in the 2--8\\,keV band \\citep{Petre02}, far less than the hard X-ray flux in the region C ($\\sim1.37\\times10^{-12}$\\,ergs\\,cm$^{-2}$\\,s$^{-1}$). We also simulated possible contamination of the strayed flux from the PWN and PSR\\,B1853$+$01 to the region A by using \\textit{xissim} \\citep{Ishisaki07}, and confirmed it is negligible. We thus conclude that the hard X-ray emission we discovered is independent from the PWN or PSR\\,B1853$+$01.   \\begin{figure}[!t] \\begin{center} \\FigureFile(70mm,70mm){figure5.eps} \\end{center} \\caption{$^{12}$CO ($J=2$--1) contours of NANTEN2 overlaid with the hard X-ray image in the 5.0--7.0\\,keV band. The blue and yellow contours are the  $^{12}$CO fluxes in the velocity ranges 40--50\\,km\\,s$^{-1}$ and 50--60\\,km\\,s$^{-1}$, respectively. The white square is the XIS FoV.}\\label{fig:co} \\end{figure}   The hard X-rays was found to have an arc-like structure (figure \\ref{fig:image}-right) and are correlated with a filament of the radio continuum (1442.5\\,MHz with VLA). For more detailed study, we compared the spatial distributions of the hard X-rays with ambient molecular clouds, $^{12}$CO ($J=2$--1) emission obtained by NANTEN2. The result is shown in figure \\ref{fig:co}. We see a clear anti-correlation between the hard X-rays and the CO cloud in the velocity range 50--60\\,km\\,s$^{-1}$ (the yellow contours). This cloud is identical to what was observed by Nobeyama 45-m radio telescope, C11 named by \\citet{Seta04}, and was suggested to be interacting with the SNRs blast wave  (see figure 11 in their paper).  The arc-like structure of the hard X-rays suggests that this emission comes from the shock. A typical shock velocity of W44 is $\\sim650\\,(kT_{\\rm e}/0.5\\,\\rm{keV})^{0.5}$\\,km\\,s$^{-1}$. As far as we assume a simple direct diffusive shock acceleration, such slow shock is insufficient to explain the flux of the observed hard X-rays. We presume that the flux enhancement possibly arises in association with the shock-cloud interaction. We focus on pre-existing cosmic ray (CR) electrons compressed and reaccelerated between the shock and dense cloud. In the radio band, the resultant synchrotron enhancement is known as the van der Laan mechanism (\\cite{Laan62a}; 1962b) which accounts for the radiative radio shell \\citep{Blandford82}. Based on a multi-wavelength observation of W44, \\citet{Cox99} pointed out that the radio synchrotron emission is possibly enhanced by the pre-existing CR electrons which are reaccelerated by the van der Laan mechanism. \\citet{Uchiyama10} applied a similar mechanism to some middle-aged SNRs including W44 to interpret their GeV $\\gamma$-ray emission. They concluded that it is sufficient for the reaccelerating of the pre-existing CR through the hadronic process to explain the observed GeV $\\gamma$-ray intensity. We simply expanded their argument to explain the observed hard X-rays. By our rough simulation, if the hard X-ray is derived from such mechanism, the electrons shall be accelerated to at least $\\sim$100\\,TeV. In this case, the acceleration time $t_{\\rm acc}$ is calculated to be \\begin{eqnarray} t_{\\rm{acc}} \\simeq 1\\times10^6 \\left(\\frac{cp_{\\rm max}}{\\rm100\\,TeV}\\right)\\left(\\frac{\\eta}{10}\\right)\\left(\\frac{v_{\\rm s}}{\\rm 100\\,km\\,s^{-1}}\\right)^{-2} \\nonumber \\\\ \\times \\left(\\frac{B_{\\rm 0}}{\\rm 25\\, \\mu \\rm{G}}\\right)^{-1} \\rm{yr}, \\label{eq:t_acc} \\end{eqnarray} where $cp_{\\rm max}$ is the maximum attainable energy, $\\eta$ is the gyrofactor \\citep{Uchiyama07}, $v_{\\rm s}$ and $B_{\\rm 0}$ are the shock velocity decelerated by cloud and the pre-shock magnetic field \\citep{Reach05}, respectively. These parameters were set as \\citet{Uchiyama10} had assumed for W44 in their calculation. As a result, the $t_{\\rm{acc}}$ becomes significantly longer than the age.  A local anti-correlation between CO clouds and non-thermal X-ray emissions is also found in another SNR, RX\\,J1713.7$-$3946 \\citep{Sano10}. Performing magnetohydrodynamic (MHD) simulations of a SNR shock passing through clumpy interstellar clouds, \\citet{Inoue12} successfully interpreted this result; the synchrotron X-ray enhancement occurs in spatially small regions at the vicinity of the cloud due to an amplified magnetic field caused by turbulent dynamo action (see also, \\cite{Giacalone07}). One major difference between RX\\,J1713.7$-$3946 and W44 is a shock velocity. \\citet{Inoue10} showed that a slow shock wave ($\\sim500$\\,km\\,s$^{-1}$) attains a maximum magnetic field strength of  $\\sim400$\\,$\\mu$G. Given a typical shock velocity of W44, $\\sim650$\\,km\\,s$^{-1}$, however, the resultant $t_{\\rm acc}$ becomes substantially longer than the synchrotron cooling time $t_{\\rm syn}$; $t_{\\rm syn}/t_{\\rm acc}<0.07$. We concluded that this scenario is also difficult to explain the observed X-ray flux.  In either case, the acceleration time $t_{\\rm acc}$ is too long to obtain TeV electrons. Alternatively, if TeV protons are still sufficiently remaining in W44, the hadronic process provides secondary electrons which may peak at TeV energies. As already claimed by \\citet{Abdo10} and \\citet{Giuliani11}, the GeV $\\gamma$-ray in W44 is possibly due to the hadronic process. We found no spatial correlation between the hard X-ray emitting region and the global GeV $\\gamma$-ray distribution (\\cite{Abdo10}; \\cite{Giuliani11}). The reason would be that the synchrotron X-ray is enhanced only around a shocked clumpy cloud where the magnetic field is amplified by the mechanism of \\citet{Inoue10}. The remaining problem is a CR escape which determines a peak energy of the existing protons, hence the secondary electrons in W44. While the maximum possible energy differs by more than one order of magnitude among theories (e.g., \\cite{Ptsukin03}), it is generally difficult for a middle-aged SNR to have 100-TeV order of protons. The origin of the hard X-rays remains an open question."
    },
    "1208/1208.6031.txt": {
        "abstract": "We investigate spectroscopically measured Ly$\\alpha$ equivalent widths and escape fractions of $244$ sources of which $95$  are Lyman Break Galaxies (LBGs) and $106$ Lyman Alpha Emitters (LAEs) at $z\\sim4.2$, $z\\sim4.8$, and $z\\sim5.6$ selected from intermediate and narrow-band observations. The sources were selected from the Cosmic Evolution Survey (COSMOS),  and observed with the DEIMOS spectrograph. We find that the distribution of equivalent widths shows no evolution with redshift for both the LBG selected sources and the intermediate/narrow-band LAEs. We also find that the Ly$\\alpha$ escape fraction of intermediate/narrow band LAEs is on average higher and has a larger variation than the escape fraction of LBG  selected sources. The escape fraction does not show a dependence with redshift. Similar to what has been found for LAEs at low redshifts, the sources with the highest extinctions show the lowest escape fractions. The range of escape fractions increases with decreasing extinction. This is evidence that the dust extinction is the most important factor affecting the escape of Ly$\\alpha$ photons, but at low extinctions other factors such as HI covering fraction and gas kinematics can be just as effective at inhibiting the escape of Ly$\\alpha$ photons. ",
        "introduction": "The study of the high redshift universe and the early evolution of galaxies has  primarily relied on two techniques to obtain large samples of high redshift galaxies, the Lyman-break technique (LBGs; Steidel et al. 1999, Ouchi et al. 2004;  Bouwens \\& Illingworth 2006, and references therein) and narrow band surveys targeting Ly$\\alpha$ emitting galaxies (LAEs;  Hu \\& McMahon 1996; Rhoads \\& Malhotra 2001; Ajiki et al. 2003; Hu et al. 2004; Taniguchi et al. 2005; Murayama et al. 2007, Gronwal et al. 2007, Ouchi et al. 2008; Hu et al. 2010, and references therein). Studying the difference in the nature and properties of the two populations, selected by these two techniques, helps to understand early stages of galaxy formation and provides constraints on reionization. However, the two populations of  galaxies are found to have a degree of overlap, with a fraction of the LBGs having Ly$\\alpha$ emission (Shapley et al. 2003; Kornei et al. 2010;  Stark et al. 2010). The varying degree of overlap between the two techniques and how it changes with redshift is still an open question. Several authors have explored this by comparing spectral energy distributions (SED) properties of these two populations (Gawiser et al. 2007; Gronwall et al. 2007). Even less understood is the degree of overlap in the Ly$\\alpha$ properties of the populations selected by these two techniques. \\citet{kornei} recently studied the Ly$\\alpha$ properties of $z\\sim3$ LBGs and found that LBGs with strong Ly$\\alpha$ emission are older, have lower SFR, and are less dusty than objects with either weak Ly$\\alpha$ emission, or the line in absorption.  They concluded that, within the LBG sample, objects with strong Ly$\\alpha$ emission represent a later stage of galaxy evolution in which supernovae-induced out\u00dfows have reduced the dust covering fraction. In contrast, analysis of LAEs at $z\\sim3.1$, $3.7$, and $5.7$ by Ouchi et al. (2008) have revealed that LAEs have lower extinction and/or younger ages than LBGs.  Due to the complex physics of Ly$\\alpha$ radiative transfer process in galaxies, modeling Ly$\\alpha$ emission, absorption, and escape has been investigated by numerous authors.  \\citet{neufeld} and \\citet{cf00}  modeled the Ly$\\alpha$ radiative transfer and investigated the role of a clumpy, dusty, multiphase ISM on Ly$\\alpha$ escape. \\citet{hansen} has expanded on these past attempts by considering the effects of several different geometrical distributions of dust clouds, while \\citet{dijkstra} and \\citet{verhammeA} have incorporated the effect of in-falling or outgoing spherical halos of neutral gas on Ly$\\alpha$ escape and its profile. In particular, the monte-carlo radiative transfer models by \\citet{verhammeB} taking into account dust, ISM kinematics, HI column densities, and gas temperature, have been able to reproduce the Ly$\\alpha$  profiles of $11$ LAEs  found in \\citet{tapken}.  Analysis of nearby Ly$\\alpha$ emitting galaxies (Kunth et al. 2003;  Mas-Hesse et al. 2003; Hayes et al. 2005; Ostlin et al. 2009; Atek et al. 2009; Scarlata et al. 2010) indicates that Ly$\\alpha$ emission is affected by ISM geometry, gas kinematics and dust. However, the order of importance of each of these factors is not clearly established and could possibly vary from object to object \\citep{schaerer}. One method to ascertain the principle physical factors that affect the Ly$\\alpha$ radiative transfer in galaxies, is to measure the Ly$\\alpha$ escape fraction (f$_{esc}$), defined as the ratio of the  observed Ly$\\alpha$ flux to what is expected from the star formation rate (SFR) of the galaxy. In recent years the study of the escape fraction of Ly$\\alpha$ photons in star forming galaxies at redshifts ranging from $z\\sim0.1-6$ has been studied by several authors (Scarlata et al. 2009; Finklestein et al. 2009; Atek et al. 2009; Hayes et al. 2010; Ono et al. 2010a; Ono et al. 2010b). Each study has found a strong trend of decreasing escape fraction with increasing extinction, though any change in the mean escape fraction of Ly$\\alpha$ sources with redshift is uncertain given the difference in the methods of selecting samples of Ly$\\alpha$ sources at  $z\\sim0.1$, $z\\sim2$, and  $z>3$.  In order to examine the varying degree of overlap between the Ly$\\alpha$ properties of these two populations (LBGs and narrow band selected LAEs) and its redshift dependence, deep spectroscopic observations are required to measure the fraction of LBGs with Ly$\\alpha$ emission. Spectroscopic follow-up for these high redshift sources has only recently been made possible  due to the technical difficulties in the spectroscopy of faint, m$_{I} >22$, high redshift sources. Ouchi et al. 2008 obtained  Subaru/FOCAS and VLT/VIMOS spectroscopy of 84 out of 858 narrow band LAE candidates at  $z= 3.1$, $3.7$, and $5.7$. The Ly$\\alpha$ luminosity function of these sources increases with redshift, indicating that galaxies with Ly$\\alpha$ emission are more common at higher redshifts. \\citet{hu2010} presented an atlas of $88$ $z\\sim5.7$ and $30$ $z\\sim6.5$ spectroscopically confirmed LAEs. Ouchi et al. 2010 presented spectra of LAEs at $z\\sim6.6$ examining the Ly$\\alpha$ line profiles, the luminosity function, clustering properties of the sources. Analysis of their sample in comparison with LAEs at z$\\sim5.7$ indicates that the intergalactic medium (IGM) was not highly neutral at $z\\sim6.6$ and the bulk of reionization of the universe occurred at $z>7$. \\citet{stark2010} spectroscopically confirmed $199$ Ly$\\alpha$ galaxies from a sample of $627$ continuum selected LBGs at $3<z<7$ and found that the fraction of LBGs with Ly$\\alpha$ emission increases with redshift and is inversely correlated with UV luminosity. The likely cause of this is a decrease in dust extinction with redshift, and also a lower HI covering fraction for sources with lower UV luminosity.   In this paper we study Ly$\\alpha$ emission from sources at  $4<z<6$, detected in deep spectroscopic survey of the COSMOS field \\citep{scoville}. The selected sources  consist of intermediate and narrow band LAEs at $z\\sim4.2$ (IA624), $z\\sim4.8$ (NB711) and  $z\\sim5.7$ (NB816), B$_J$ LBGs, g$^+$ LBGs, V$_J$ LBGs, r$^+$ LBGs, i$^+$ LBGs, and sources with  photometric redshifts $z > 4$ . In \\S2, we present the data, and the method used for source selection. In \\S3, we present our analysis of the Ly$\\alpha$ emission as it relates to both redshift and our source selection. In \\S4 we estimate the Ly$\\alpha$ escape fraction  and perform a speculative analysis based on our estimates. Our conclusions are presented in \\S5. We assume  {\\it H$_o$}$=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.3$, and  $\\Omega_{\\Lambda}=0.7$. We assume AB magnitude. s ",
        "conclusions": "In this paper we present an analysis of  a spectroscopic sample of $244$ LBGs and LAEs at $4<z<6$ in COSMOS with clear Ly$\\alpha$ detections. We have attempted to determine variations in the Ly$\\alpha$ properties for these sources and their evolution with redshift. The sources were targeted for spectroscopy using a range of high redshift selection techniques, including LBG, intermediate/narrow-band, photo-z, and IRAC CH2 detections. The goal of the spectroscopic program was to select as complete a sample at $z>4$ as possible, for objects brighter than $z^+<25$ and more massive than $10^{10.5}M_\\sun$ (Capak et al. in prep). We  measured EW$_{Ly\\alpha,0}$ and escape fractions for $B_J$,$g^+$,$V_J$,$r^+$,$i^+$ LBGs, one intermediate-band and two narrow-band selected samples of LAEs at $z\\sim4.2$, $z\\sim4.8$, and $z\\sim5.6$. A sub-sample of $153$ sources have estimates of E(B-V), SFR and M$_\\sun$ from SED modeling. We analyze the variations of the Ly$\\alpha$ properties for this subset with respect to these parameterizations of the host galaxies. The results are summarized below.  1) We find that the Ly$\\alpha$ EWs remain roughly constant with redshift for both the LBG  and intermediate/narrow-band LAEs. While low EW$_{Ly\\alpha,0}$ are detected for sources at all redshifts, increasingly larger EW$_{Ly\\alpha,0}$ are measured for sources from samples at higher redshifts. These results are in accordance with the results of Stark et al. 2010 who found a similar trend for LBGs with Ly$\\alpha$ at $z=3-6$, and with the similar findings of \\citet{nilsson} studying LAEs at lower redshifts (z=$2-3$). The speculation is that the change in EW distributions with redshift is the result of increased dust content  in LAEs at lower redshifts, but this is yet to be confirmed.  2) No trends were found between Ly$\\alpha$ luminosity  and stellar mass or SFR. Except for the IA624 LAEs, which on average have lower UV luminosities, the sources tend to have similar stellar masses and SFRs. The mean Ly$\\alpha$ luminosities are slightly higher for the LAEs than the LBGs.  3) We find that the Ly$\\alpha$ escape fraction of narrow-band LAEs is, on average, higher and has a larger variation than LBG selected sources. The escape fraction does not show a dependence on redshift.  Our escape fraction for NB816 LAEs, 0.48, agrees within the errors to escape fraction of  NB816 selected sources measured by \\citet{onoB} in the Subaru/XMM-Newton Deep Survey field (0.36), and the mean escape fraction of  Ly$\\alpha$ sources (0.32) at $z=2.2$ studied by  Hayes et al. (2010).   4) Similar to what has been found for sources with Ly$\\alpha$ emission at low redshifts, the sources with the highest extinctions show the lowest escape fractions. The range of escape fractions increases with decreasing extinction. This is evidence that the dust extinction is the most important factor affecting the escape of Ly$\\alpha$ photons, but at low extinctions other factors such as HI covering fraction and gas kinematics can be just as effective at inhibiting the escape of Ly$\\alpha$ photons."
    },
    "1208/1208.1265_arXiv.txt": {
        "abstract": "{} {We investigate the influence of stellar bulges on the star formation and morphology of disc galaxies that suffer from ram pressure. Several tree-SPH (smoothed particle hydrodynamics) simulations have been carried out to study the dependence of the star formation rate on the mass and size of a stellar bulge. In addition, different strengths of ram pressure and different alignments of the disc with respect to the intra-cluster medium (ICM) are applied.} {The simulations were carried out with the combined N-body/hydrodynamic code GADGET-2 with radiative cooling and a recipe for star formation. The same galaxy with different bulge sizes was used to accomplish 31 simulations with varying inclination angles and surrounding gas densities of $10^{-27}\\, \\mathrm{g}\\,\\mathrm{cm}^{-3}$ and $10^{-28}\\,\\mathrm{g}\\,\\mathrm{cm}^{-3}$. For all the simulations a relative velocity of $1000\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$ for the galaxies and an initial gas temperature for the ICM of $10^7\\,\\mathrm{K}$ were applied. Besides galaxies flying edge-on and face-on through the surrounding gas, various disc tilt angles in between were used. To allow a comparison, the galaxies with the different bulges were also evolved in isolation to contrast the star formation rates. Furthermore, the influence of different disc gas mass fractions has been investigated.} {As claimed in previous works, when ram pressure is acting on a galaxy, the star formation rate (SFR) is enhanced and rises up to four times with increasing ICM density compared to galaxies that evolve in isolation. However, a bulge suppresses the SFR when the same ram pressure is applied. Consequently, fewer new stars are formed because the SFR can be lowered by up to $2\\,\\solarmass\\, \\mathrm{yr}^{-1}$. Furthermore, the denser the surrounding gas, the more inter-stellar medium (ISM) is stripped. While at an ICM density of $10^{-28}\\,\\mathrm{g}\\,\\mathrm{cm}^{-3}$ about 30\\% of the ISM is stripped, the galaxy is almost completely (more than 90\\%) stripped when an ICM density of $10^{-27}\\mathrm{g}\\,\\mathrm{cm}^{-3}$ is applied. But again, a bulge prevents the stripping of the ISM and reduces the amount being stripped by up to 10\\%. Thereby, fewer stars are formed in the wake if the galaxy contains a bulge. The dependence of the SFR on the disc tilt angle is not very pronounced. Hereby a slight trend of decreasing star formation with increasing inclination angle can be determined. Furthermore, with increasing disc tilt angles, less gas is stripped and therefore fewer stars are formed in the wake. Reducing the disc gas mass fraction results in a lower SFR when the galaxies evolve in vacuum. On the other hand, the enhancement of the SFR in case of acting ram pressure is less pronounced with increasing gas mass fraction. Moreover, the fractional amount of stripped gas does not depend on the gas mass fraction.} {} ",
        "introduction": "As is known since the 1970s, the star formation rate of galaxies depends heavily on their environment. As stated by \\citet{bo78}, cluster galaxies are found to be redder than field galaxies. This implies that cluster galaxies must form stars at a lower rate than their counterparts in the field. Although the star formation rate is lower in cluster galaxies, they also tend to have less gas than galaxies in less dense environments. \\\\ Furthermore, an excess of blue objects has been found at redshift $z = 0.5$, compared to lower redshifts. Moreover, spiral galaxies have a higher star formation rate than ellipticals. Ellipticals on the other hand are more numerous in the centre of a galaxy cluster \\citep{bark09}. Furthermore, \\citet{dressler97} found that between redshift $z = 0.5$ and $z = 0$ the fraction of S0 galaxies is strongly increasing. Hence, much more early-type galaxies can be found in clusters, which could be explained by different galaxy transformation processes in clusters of galaxies. \\\\ \\noindent Many different mechanisms have been studied in this context to explain the change in morphology and star formation of cluster galaxies. Among others, galaxy harassment \\citep[e.g.][]{moore98} and strangulation/starvation \\citep[e.g.][]{larson80} are discussed mechanisms. Furthermore, galaxy-galaxy interactions \\citep[e.g.][]{sulentic76, bushouse87, sosa88, kewley06} as well as ram-pressure stripping \\citep{gunn_gott72} can be found. Because these different processes can happen at the same time, it is crucial to separate these processes which can be done by studying the evolution of the star formation rate in dependence on the environment \\citep[e.g.][]{poggianti06}. \\\\ \\noindent Clusters of galaxies contain a hot ($\\approx 10^7\\,\\mathrm{K}$) and very thin ($10^{-29} - 10^{-26} \\, \\mathrm{g}\\,\\mathrm{cm}^{-3}$) gas, the so-called intra-cluster medium, detected by X-ray observations. Galaxies that are moving in the cluster's potential and consequently are exposed to the ICM, are feeling the ram pressure of this hot gas and the inter-stellar medium is stripped if the ram pressure exceeds the gravitational restoring force, according to the Gunn and Gott criterion \\citep{gunn_gott72}. Additionally, ram-pressure stripping is also a mechanism for enriching the ICM with metals \\citep[e.g.][]{schindler08, lovisari09, durret11}. The influence of ram-pressure stripping on the morphology and star formation was exhaustively investigated in the last decade \\citep[e.g.][]{gavazzi01, yoshida04, scott10, kenney04}, by investigating HI deficiencies and $\\mathrm{H}\\alpha$ and CO observations to investigate regions of star formation. This showed that interactions with the ICM cause the galaxies to have morphological disturbances of the gas disc and extensive gas loss in the central part of a galaxy cluster \\citep[e.g.][]{boselli09}. In addition, star formation has also been found outside the galaxies themselves \\citep[e.g.][]{hester10}. \\\\ \\noindent Another approach to study the different interaction processes of galaxies in a cluster are numerical simulations. Tidal interactions between galaxies and galaxy-galaxy mergers were investigated by \\citet{kronberger06, kronberger07}. Additionally, they provided studies of the influence on the rotation curves and velocity fields of spiral galaxies \\citep{kronberger07, kapferer06}. Furthermore, it has been found that nonaxisymmetric and nonbisymmetric features are introduced by tidal interactions \\citep{rubin99}. Ram-pressure stripping simulations were also performed by several groups, which found a good agreement of the Gunn \\& Gott criterion in hydrodynamic simulations as well \\citep{abadi99, vollmer01, roediger05}. \\citet{jachym09, jachym07} investigated the influence of a time-varying ram pressure and different inclination angles of the galaxies relative to the movement direction of a galaxy. These simulations were performed mostly with a smoothed particle hydrodynamics (SPH) prescription. Other methods, including sticky particles \\citep[e.g.][]{vollmer01} and Eulerian grid-codes \\citep[e.g.][]{roediger06, roediger07} were also used to simulate ram-pressure stripping. \\\\ \\noindent Because one of the main aims of this paper is to study the evolution of the star formation rate depending on the environment, the simulation code includes a star formation recipe. A high pressure exerted by the ICM to the ISM can enhance the star formation rate significantly, as e.g. found by \\citet{kronberger08}. Other simulations also confirm the enhancement of the star formation rate when ram pressure is applied. \\citet{bekki03} demonstrated that the pressure of the ICM can induce the collapse of molecular clouds and consequently trigger a burst of star formation. Also using SPH, their recipe for star formation employs the local dynamical timescale and the sound crossing time of the gas particles to transform them into stellar particles. Furthermore, \\citet{bekki11} used a Schmidt-law with exponent $\\gamma = 1.5$ to control the rate at which new stars are produced. \\citet{vollmer12} also use a model for star formation in their simulation prescription, adopting the sticky-particles approach for the hydrodynamics with ram pressure applied by an additional acceleration to the particles. Stars are being formed during cloud collisions, represented by tracer-particles with zero mass. These particles are then evolved passively. \\\\ In our approach, star formation processes are not restricted to the galaxy or some parts of the galaxy itself but are separately calculated for each particle, considering temperature and density in a self-consistent way throughout the whole simulation domain. In the same way, new stellar particles are spawned that conserve the phase-space properties of the gas particle, containing the mass of newly formed stars. Thus, enhancements in the star formation rate due to ram pressure are obtained self-consistently. Already adopted for ram-pressure stripping simulations in \\citet{kronberger08}, we describe this recipe in more detail in section \\ref{sim_setup}. \\\\ \\noindent In this work, especially the influence of the combination of a galaxy's morphology and ram-pressure stripping on the star formation rate is investigated. In this context, the special case of the influence of a stellar bulge is studied. We use different strengths of ram pressure and different model galaxies to decode a bulge's influence on ram-pressure stripping. The paper is organised as follows: in Sect. 2 we present the numerical setup. In Sect. 3 the results are presented and are compared to observations in Sect. 4. Finally, in Sect. 5 a summary with the main conclusions is given. \\\\ ",
        "conclusions": "We carried out numerical studies on the effect of ram-pressure stripping and especially the influence of a galaxy's bulge. Thirty-one different high-resolution N-body/SPH simulations were conducted. The size of the galaxy bulge was varied and we used two different ICM densities of $10^{-28}\\,\\mathrm{g}\\,\\mathrm{cm}^{-3}$ and $10^{-27}\\,\\mathrm{g}\\,\\mathrm{cm}^{-3}$. In each case, the relative velocity of the galaxies was $1000\\,\\mathrm{km}\\,\\mathrm{s}^{-1}$ and the temperature of the ICM was $10^7$ K. \\\\  - The size and mass of the bulge of the model galaxies were varied. While the total mass of all galaxies remained the same, mass was shifted from the dark-matter halo to the bulge in the initial conditions. Hence, a deeper potential well was formed. This induced a higher gas mass concentration in the centre of the galaxy, which increased the star formation. \\\\  - After 2 Gyr of evolution in isolation, the spiral structure of the ISM developed into concentric rings in the simulation. When ram pressure was acting on a galaxy, however, the ISM was compressed, which by itself caused an enhancement of the star formation rate. Moreover, the ISM was stripped and the galaxy can lose most of its gas up to 90\\% if there was high ram pressure. In this case of high ram pressure acting, the process was basically completed within 200 Myr. Owing to the steeper potential well of the same galaxy containing a bulge, less gas was stripped as stated by the Gunn \\& Gott criterion \\citep{gunn_gott72}. As confirmed by previous investigations, the inclination angle does not have a big influence on the amount of stripped gas. Only if the inclination was nearly edge-on the fraction of gas in the wake of the galaxy was much lower. A difference in the amount of stripped gas for distinct inclination angles can be found when lower ram pressure is acting on the galaxies. \\\\  - Star formation was enhanced when ram pressure was acting on a galaxy. In the considered ram-pressure scenarios, the star formation rate was up to four times higher than for the same galaxy evolving in vacuum. The higher the ram pressure, the more gas was stripped and therefore fewer stars were formed in the disc. On the other hand, many more dense star forming gas knots were present in the wake, and the higher the external pressure on these knots, the more stars were formed. \\\\  - When the inclination angle of a galaxy was modified, the star formation rate did not change significantly when low ram pressure acted on a galaxy. Still, the highest SFR could be found when the galaxy was flying face-on through the ICM. In the high-density ICM, increasing the inclination angle resulted in a slight decrease of the star formation rate. Furthermore, in the edge-on case the star formation rate was lower at the beginning, but higher at the end of the simulation. \\\\  - The mass fraction of newly formed stars in the wake compared to the total mass was much higher for stronger ram pressure. Again, because less gas was stripped if the bulge mass was higher, fewer stars were formed in the wake when a big bulge was present. In agreement with previous results, less mass of newly formed stars was present in the wake after 500 Myr of evolution if the bulge size of a galaxy was increased, and obviously also the mass in the wake was higher when the ram pressure was increased. \\\\  - Different disc gas mass fractions do have a direct influence on the star formation rate, which decreases with a lower disc gas mass fraction. Interestingly, the enhancement due to ram-pressure was less pronounced for larger gas mass fractions. Apparently, because the gas is less dense when the gas mass fraction is low, ram-pressure can compress the gas to bring it into a star-forming regime. On the other hand, when the gas mass fraction and hence the density was already high, the compression effect was smaller. \\\\  - Our investigations show that for high ram pressure up to 25\\% of the newly formed stars were formed in the wake of the galaxy, whereas with low ram pressure, only a marginal amount of around 1\\% of all newly formed stars were present in the wake. When the ICM density was high, gas knots were formed from the stripped material in the wake. These high dense structures contain a cool core and hence many new stars were formed in these self-gravitating gas knots. With evolving time the gas knots are losing mass due to star formation processes. Although the gas knots were heated by stellar feedback, and gas flows out via stellar winds, the gas knots remained intact. Obviously, the cooling was very efficient. Additionally, self-gravity and the external pressure from the ICM kept the gas knots dense. Hence, the mass loss was almost completely due to star formation. The stars in these knots were not subject to ram pressure and were attracted by the galaxy itself. Because the stars are collisionless, they fall through the disc and back again, oscillating. Therefore, stars were present also in front of the disc. \\\\  - The presence of these gas knots corresponds well to observations of the 'fireballs' in the wake of RB199 in the Coma cluster by \\citet{yoshida08} and the dense blue gas knots in the wake of IC 3418 in the Virgo cluster \\citep{martin05}. In both cases, the most plausible mechanism to form these gas knots is ram-pressure stripping, also confirmed by the simulations we performed here. \\\\"
    },
    "1208/1208.3260_arXiv.txt": {
        "abstract": "We present results from Keck/NIRSPEC and Magellan/MMIRS follow-up spectroscopy of Ly$\\alpha$ emitters (LAEs) at $z=2.2$ identified in our Subaru narrowband survey. We successfully detect \\Ha\\ emission from seven LAEs, and perform a detailed analysis of six LAEs free from AGN activity, two out of which, CDFS-3865 and COSMOS-30679, have \\OII\\ and \\OIII\\ line detections. They are the first \\OII-detected LAEs at high-$z$, and their \\OIII/\\OII\\ ratios and $R23$-indices provide the first simultaneous determinations of ionization parameter and oxygen abundance for LAEs. CDFS-3865 has a very high ionization parameter ($q_{ion}=2.5^{+1.7}_{-0.8} \\times 10^8$\\,cm\\,s$^{-1}$) and a low oxygen abundance ($12+\\log ({\\rm O/H})=7.84^{+0.24}_{-0.25}$) in contrast with moderate values of other high-$z$ galaxies such as Lyman-break galaxies (LBGs). COSMOS-30679 also possesses a relatively high ionization parameter ($q_{ion}=8^{+10}_{-4} \\times 10^7$\\,cm\\,s$^{-1}$) and a low oxygen abundance ($12+\\log ({\\rm O/H})=8.18^{+0.28}_{-0.28}$). Both LAEs appear to fall below the mass-metallicity relation of $z\\sim 2$ LBGs. Similarly, a low metallicity of $12+\\log ({\\rm O/H})<8.4$ is independently indicated for typical LAEs from a composite spectrum and the \\NII/\\Ha\\ index. Such high ionization parameters and low oxygen abundances can be found in local star-forming galaxies, but this extreme local population occupies only $\\sim 0.06$\\% of the SDSS spectroscopic galaxy sample with a number density $\\sim 100$ times smaller than that of LAEs. With their high ionization parameters and low oxygen abundances, LAEs would represent an early stage of galaxy formation dominated by massive stars in compact star-forming regions. High-$q_{ion}$ galaxies like LAEs would produce ionizing photons efficiently with a high escape fraction achieved by density-bounded \\HII\\ regions, which would significantly contribute to cosmic reionization at $z>6$. ",
        "introduction": "\\label{sec:introduction}   \\Lya\\ emitters (LAEs), galaxies commonly observed at high redshifts with strong \\Lya\\ emission, are considered to be low-mass, young galaxies as suggested from their small sizes, faint continua, and low masses inferred from spectral energy distribution (SED) fitting (e.g., \\citealt{venemans2005,gawiser2006,pirzkal2007,overzier2008,malhotra2012}). LAEs are therefore likely to represent galaxies in the early stages of galaxy evolution. More directly, \\citet{cowie2011} investigate rest-frame optical nebular lines of LAEs at low redshifts ($z\\sim0.3$). The authors find that a large portion of LAEs ($75$\\,\\%) have equivalent width (EW) of \\Ha\\ $>100$\\,\\AA, and that LAEs on average have lower metallicities and younger ages than the UV-continuum sample. These findings are consistent with the idea that LAEs are galaxies in early stages of galaxy formation. At higher redshifts, LAEs are efficiently detected thanks to narrowband imaging techniques which have enriched our knowledge about the younger universe (e.g., \\citealt{CH1998,MR2002,ouchi2003,gawiser2006,shimasaku2006,% gawiser2007,gronwall2007,ouchi2008,nilsson2009,guaita2010,finkelstein2011,% nakajima2012}). However, more direct evidence supporting the idea that high redshift LAEs are in an early evolutionary phase of formation is still needed.   The gas-phase metallicity is a key property of galaxies, since it is a record of their star-formation histories. This physical quantity is relatively easily constrained with line ratios of nebular lines at rest-frame optical wavelengths (e.g., \\citealt{pagel1979,KD2002}, and references therein). Another key quantity is the ionization parameter, defined as the ratio of the mean ionizing photon flux to the mean hydrogen atom density. Since the excitation of the \\HII\\ region is sensitive to the age distribution of the exciting stars, the ionization parameter provides a rough estimate of age of a galaxy (e.g., \\citealt{dopita2006}). A large ionization parameter is observed from a galaxy dominated with massive stars, which is a sign that the galaxy is in an early stage of galaxy formation. In addition, the ionization parameter depends on the optical depth in an \\HII\\ region (e.g., \\citealt{brinchmann2008}); a large ionization parameter may be instead due to a low optical depth and a high escape fraction of ionizing photons. Therefore, constraining ionization parameters for LAEs may provide an independent clue to whether or not LAEs, commonly observed star-forming galaxies at high redshift, significantly contribute to the cosmic reionization in the early universe. Determining the ionization parameter requires two emission lines of different ionization stages of a same element, such as \\OII\\ and \\OIII\\ (e.g., \\citealt{KD2002}).   However, well-defined samples of LAEs are generally located at very high redshifts ($3<z<7$), where rest-frame optical nebular lines are redshifted into infrared wavelengths. Recently, \\citet{finkelstein2011} and \\citet{nakajima2012} instead take advantage of LAEs at moderately high redshifts ($z\\sim 2$), where the nebular lines up to \\Ha\\ are observable from the ground in the near-infrared (NIR) windows.   These studies extend the metallicity censuses to galaxies with strong \\Lya\\ emission. \\citet{finkelstein2011} obtain NIR spectra for two $z\\sim 2.4$ LAEs, and find that at least one LAE appears to be less chemically enriched than $z\\sim 2$ continuum-selected galaxies at similar stellar masses \\citep{erb2006a}. \\citet{nakajima2012} estimate average line fluxes for \\OII$\\lambda 3727$ and \\Ha+\\NII$\\lambda6584$ for $z=2.2$ LAEs by stacking $1.18$ and $2.09\\,\\mu$m narrowband images for more than $100$ LAEs, and place a firm lower-limit for the average metallicity of this population. Interestingly, the lower-limit is higher than expected for its stellar mass at the redshift.   In contrast to the gas-phase metallicities, the ionization parameters of LAEs are unknown. Such measurements have only been obtained for bright galaxies such as Lyman-break galaxies (LBGs) or strongly lensed galaxies at high-$z$ (e.g., \\citealt{pettini2001,hainline2009,richard2011}). Although the current number of measurements is small, ionization parameters of such high-$z$ galaxies are likely to be higher on average than those of local galaxies. A comparison of ionization parameters for LBGs and LAEs will enable us to discuss whether or not LAEs are young galaxies at high redshifts.   In order to obtain reliable measurements of properties including metallicity and ionization parameter of LAEs at high redshifts, we have carried out a large survey for $z=2.2$ LAEs, using our custom narrowband filter NB387 with Subaru/Suprime-Cam. At $z=2.2$, important nebular lines such as \\OII$\\lambda 3727$, \\Hb, \\OIII$\\lambda\\lambda 5007,4959$, \\Ha, \\NII$\\lambda 6584$, are observable from the ground. Initial results were based on using three narrowbands to detect \\Lya, \\OII, and \\Ha\\ over the same volumes (\\citealt{nakajima2012}; see also \\citealt{lee2012} and \\citealt{ly2011} for the two NIR narrowbands and the New\\Ha\\ Survey).   In this paper, we present results from NIR spectroscopy. We used Keck/NIRSPEC and Magellan/MMIRS spectrographs, and successfully detected \\Ha\\ emission from seven LAEs. The number is double the previous number of high-$z$ LAEs with NIR spectra (two from \\citealt{finkelstein2011} and two from \\citealt{mclinden2011}), and allows us to begin to examine statistical variation of rest-frame optical spectroscopic properties of LAEs at $z\\sim 2$. Our first NIR spectroscopic result discusses the kinematics of LAEs and is presented in \\citet{hashimoto2013}. As a companion study, this paper presents the ionization and chemical properties of LAEs based on multiple nebular lines. Remarkably, we detected \\OII\\ and \\OIII\\ lines from two LAEs , which provide ionization parameter estimates for LAEs for the first time. In addition, the oxygen lines allow us to determine oxygen abundances that compliment the previous spectroscopic constraints on metallicity of LAEs from the \\NII/\\Ha\\ index.   We also investigate \\Lya\\ and \\Ha\\ hydrogen lines for LAEs. A comparison of observed \\Lya/\\Ha\\ ratios with the Case B recombination value (\\Lya/\\Ha\\ $=8.7$; \\citealt{brocklehurst1971}) provides important insights into the physical mechanisms causing the strong \\Lya\\ emission in LAEs. Furthermore, we study \\Lya\\ and \\Ha\\ equivalent widths which probe star-formation history, stellar age, and metallicity.   This paper is organized as follows. We describe the data in \\S\\ref{sec:data}. The detection and measurement of emission lines in the NIR spectroscopy is summarized in \\S\\ref{sec:line_measure}. In \\S\\ref{sec:properties}, we derive properties of LAEs including estimates of ionization parameter, metallicity, and SFR from the rest-frame optical nebular lines. We also check for the presence of active galactic nuclei (AGN) in the LAEs. In \\S\\ref{sec:discussion}, we compare LAEs with other galaxies in terms of their ionization state, metallicity, and SFR. We then discuss the implications. We also discuss the physical properties inferred from \\Lya\\ and \\Ha\\ emission. We conclude the paper in \\S\\ref{sec:summary} with a summary. Throughout this paper, magnitudes are given in the AB system \\citep{oke1974}, and we assume a standard $\\Lambda$CDM cosmology with $(\\Omega_m,\\Omega_{\\Lambda},H_0)= (0.3,0.7,70\\,{\\rm km}\\,{\\rm s}^{-1}\\,{\\rm Mpc}^{-1})$. ",
        "conclusions": "\\label{sec:discussion}  In this section, we discuss the physical properties of LAEs determined by this study in conjunction with \\citet{finkelstein2011} and \\citet{nakajima2012}. In the first subsection, we make comparisons in terms of ionization parameter, metallicity, and SFR of LAEs at $z\\sim 2$ as well as lower-redshifts, and other galaxies at similar redshifts such as LBGs. We then examine the implications. In the second subsection, we show the properties of LAEs newly found by spectroscopy based on \\Lya\\ and \\Ha\\ hydrogen recombination lines. We extend the discussion to understanding the origins of the strong \\Lya\\ emission observed in LAEs.      \\subsection{Comparisons between LAEs and other galaxies} \\label{ssec:comparison}  \\subsubsection{Ionization State} \\label{sssec:high_qion}  We find high ionization parameters for the LAEs. In this section, we examine this point further by comparison to other galaxies, and discuss its implications.   The situation is most clearly seen in Figure \\ref{fig:o3o2r23}, which shows the \\OIII/\\OII\\ ratio versus $R23$-index. We plot CDFS-3865 and COSMOS-30679 (red circles), and other galaxies such as LBGs at $z=2$--$3$ (blue symbols). The black grid represents model predictions of \\OIII/\\OII\\ ratio and $R23$-index at a given metallicity and (discrete) ionization parameter \\citep{KD2002}. According to the calculation, CDFS-3865 is close to the $q_{ion}=3\\times 10^8$\\,cm\\,s$^{-1}$ curve. This confirms its high ionization parameter and low oxygen abundance. The high ionization parameter is suggestive of CDFS-3865 showing a very hard ionizing spectrum. Therefore, one of the most straightforward interpretations is that CDFS-3865 is a young galaxy dominated by massive stars in possibly small star-forming regions (expected for young \\HII\\ regions). Its low metallicity ($Z\\sim 0.1\\,Z_{\\odot}$) and large EW(\\Ha) ($\\sim 800$\\,\\AA) also support this idea (see also \\S\\ref{sssec:MZR}). The idea of LAEs being less-evolved galaxies is consistent with previous photometric results, such as low stellar masses (e.g., \\citealt{pirzkal2007,ono2010b}) and small sizes (e.g., \\citealt{bond2009,malhotra2012}).   Compared to CDFS-3865, most other $z=2$--$3$ galaxies appear to have lower ionization parameters. Notable exceptions are objects found by \\citet{fosbury2003}, \\citet{erb2010}, and (partly) \\citet{richard2011}. \\citet{fosbury2003} and \\citet{erb2010} reveal from their multi-emission lines analysis that their galaxies possess high ionization parameters ($q_{ion}\\gtrsim 10^9$\\,cm\\,s$^{-1}$), low metallicities ($Z\\sim 0.1\\,Z_{\\odot}$ or less), being roughly comparable to (or more or less extreme than) those found in CDFS-3865. Interestingly, both galaxies turn out to exhibit strong \\Lya\\ emission. This supports our suggestion that strong \\Lya\\ galaxies are in part represented by young galaxies. \\cite{richard2011} also find two lensed galaxies (blue diamonds with upward arrow) possibly have high ionization parameters ($q_{ion}\\gtrsim 10^8$\\,cm\\,s$^{-1}$) judged from their metallicities determined empirically (\\Oabundance\\ $=8.00^{+0.44}_{-0.50}$ and $8.77^{+0.14}_{-0.14}$, for Sextet and MACS0712, respectively). Unfortunately, \\Lya\\ measurements are not given in \\cite{richard2011}.   However, as indicated by COSMOS-30679, LAEs' ionization parameters are unlikely to be always very high. The ionization parameter for COSMOS-30679 is estimated to be $q_{ion}= 8^{+10}_{-4}\\times 10^7$\\,cm\\,s$^{-1}$, relatively high but comparable to those found in other high-$z$ galaxies (e.g., \\citealt{hainline2009}) or LBGs (e.g., \\citealt{mannucci2009}), although the error is very large. COSMOS-30679 is thus considered to be more evolved than CDFS-3865, suggested from its relatively higher metallicity, smaller EW(\\Ha), older age estimated from SED fitting, and smaller ionization parameter.   Next, we compare our LAEs with galaxies seen in the local universe. In Figure \\ref{fig:o3o2r23}, the gray dots indicate SDSS galaxies with \\Oabundance\\ $\\gtrsim 8.2$ (\\citealt{tremonti2004,KD2002}) whose metallicities are derived by applying photoionization models to the most prominent optical emission lines (\\OII, \\Hb, \\OIII, \\Ha, \\NII, \\SII). They thus represent typical local galaxies. According to \\citet{dopita2006}, the SDSS galaxies have ionization parameters $q_{ion}={\\rm several}\\times 10^7$\\,cm\\,s$^{-1}$ on average. Galaxies at $z=2$--$3$ including the LAEs appear to have systematically higher ionization parameters than the SDSS galaxies. Galaxies at intermediate redshift ($z\\sim 0.7$; \\citealt{lilly2003}) are also plotted in Figure \\ref{fig:o3o2r23}. They appear to fall between $z\\gtrsim 2$ galaxies and the SDSS galaxies in Figure \\ref{fig:o3o2r23}, showing systematically higher \\OIII/\\OII\\ ratios than the SDSS galaxies. Since their metallicities are comparable to those of the SDSS galaxies \\citep{lilly2003,tremonti2004}, the higher \\OIII/\\OII\\ ratios are likely due to their somewhat higher ionization parameters. Combined with the findings of much higher \\OIII/\\OII\\ ratios in $z=2$--$3$ galaxies, higher-$z$ galaxies tend to have higher ionization parameters.   Another interesting comparison is with the low-metallicity galaxies in the local universe. The gray diamonds in Figure \\ref{fig:o3o2r23} show local low-metallicity galaxies with $7\\lesssim$ \\Oabundance\\ $\\lesssim 8.5$ \\citep{nagao2006} whose metallicities are measured with the {\\it direct T$_e$} method (e.g., \\citealt{izotov2006}). Unlike typical SDSS galaxies, they appear in almost the same parameter space occupied by high-$z$ galaxies on Figure \\ref{fig:o3o2r23}. Since their low-metallicities are determined by the {\\it direct T$_e$} method and thus reliable, the high \\OIII/\\OII\\ ratios can be interpreted as due to high ionization parameters. Indeed, from the comparison between \\OIII/\\OII\\ ratios and photo-ionization models given in \\citet{nagao2006}, it is inferred that local low-metallicity galaxies typically have higher ionization parameters ($\\sim 8\\times 10^7$\\,cm\\,s$^{-1}$, or higher) than more metal-rich galaxies seen in the SDSS sample. Therefore, these low-metallicity galaxies in the local universe appear to have similar metallicities and ionization parameters as high-$z$ star-forming galaxies including LBGs, and particularly those with extreme quantities are likely analogs of high-$z$ LAEs such as CDFS-3865 in terms of their high ionization parameters and low metallicities.   One such extreme population is star-forming luminous compact galaxies (LCGs), or popularly referred to as ``green pea\" galaxies (GPs; \\citealt{cardamone2009}). The GPs are found in the SDSS spectroscopic sample% \\footnote{ Hence some of the GPs overlap with the local low-metallicity galaxies plotted with the gray diamond on Figure \\ref{fig:o3o2r23}.} to have green colors in the SDSS $gri$ composite images, caused by the very strong \\OIII\\ emission lines with EWs of $\\sim 1000$\\,\\AA\\ falling into the $r$ band. They show low metallicities ($Z\\sim 0.2\\,Z_{\\odot}$; \\citealt{amorin2010}) and high ionization parameters ($\\gtrsim 10^8$\\,cm\\,s$^{-1}$) inferred from their high \\OIII/\\OII\\ ratios \\citep{jaskot2013}, analogous to those observed in the LAEs. Their SFRs are also comparable to those for LAEs \\citep{izotov2011}, in contrast with more modest SFRs seen in typical local low-metallicity galaxies (e.g., \\citealt{lee2004}). A similar discussion is also provided by \\citet{xia2012}.   In the SDSS Date Release 7 (DR7) spectroscopic galaxy sample, $251$ GPs are found \\citep{cardamone2009}. Searching the entire SDSS DR7 spectroscopic sample yields $418,429$ objects which fall in the GPs redshift range ($z=0.112$--$0.360$; \\citealt{cardamone2009}) and are classified as GALAXY. Therefore, the GPs occupy only $\\sim 0.06$\\,\\% of the SDSS galaxy sample. Based on the luminosity function determined for the SDSS galaxies \\citep{blanton2003}, the number density is $\\sim 1.5\\times 10^{-2}$\\,Mpc$^{-3}$ above a detection limit $M_r=-17.63$, which corresponds to $0.037\\,L^{\\star}$. The limit is calculated from the apparent magnitude limit ($r<17.77$; \\citealt{tremonti2004}) for the spectroscopic sample and the redshift threshold ($z>0.028$; \\citealt{nagao2006}). The number density of the GPs is thus approximately $9\\times 10^{-6}$\\,Mpc$^{-3}$, which is almost consistent with the number density found by \\citet{cardamone2009} ($\\sim 2$ GPs per deg$^2$). On the other hand, the number density of LAEs at $z=2.2$ is $\\sim 1.7\\times 10^{-3}$\\,Mpc$^{-3}$ calculated by integrating the luminosity function provided by \\citet{hayes2010}% \\footnote{ The detection limit is $L=2.8\\times 10^{41}$\\,\\ergs\\ ($\\sim 0.02\\,L^{\\star}$).} above $L=0.037\\,L^{\\star}$. The number density of the GPs is thus almost two order of magnitude smaller than that of LAEs. Therefore, if the GPs alone represent galaxies with high ionization parameter in the local universe, although it is still uncertain, the large difference of abundances between LAEs and the GPs suggests that such galaxies appear to be much more abundant at high-$z$ while rarely seen in the local universe.   The final noteworthy implication of the high ionization parameters found in high-$z$ LAEs is that such galaxies may provide additional photons that cause hydrogen reionization of the IGM in the early universe. Previous censuses of early galaxies have revealed a possible shortage of ionizing photons for the cosmic reionization (e.g., \\citealt{ouchi2009,robertson2010}). \\citet{ouchi2009} find that the universe could not be totally ionized by only galaxies at $z=7$ if there is no evolution of properties (e.g., escape fraction of ionizing photons (\\fesc), metallicity, dust extinction) from $z\\sim 3$ to $z=7$. One approach to resolve the problem is to adopt larger \\fesc, which may be the case for galaxies with higher ionization parameter. Since lower ionization species such as \\OII\\ dominate in the outer regions of an usual, ionization-bounded \\HII\\ region (e.g., \\citealt{shields1990,oey1997}), a density-bounded \\HII\\ region will show a larger \\OIII/\\OII\\ ratio. Therefore, an ionization parameter depends not only on the hardness of the ionization radiation field but on the optical depth in an \\HII\\ region (e.g., \\citealt{brinchmann2008, giammanco2005}); higher ionization parameters are expected to be observed from optically thin, density-bounded \\HII\\ regions. E.g., \\citet{giammanco2005} calculate that in the condition of density bounding, an \\HII\\ region with \\fesc=$0.5$ shows the same ionization parameter as a region with $\\sim\\times 10$ greater ionizing flux without any photon escape. Consequently, the high ionization parameters found in the LAEs may be due to a low optical depth and a high \\fesc, probably achieved by density-bounded \\HII\\ regions.  \\begin{figure*} \\epsscale{1.15} \\plotone{Mass_Z.eps} \\caption{ (left) Mass-metallicity (M-Z) relation for $z\\sim 2$ LAEs (red symbols), lensed galaxies, and continuum selected galaxies (blue). The circles show our LAEs, whose metallicities are measured/constrained by using $R23$-index (with error bar)/$N2$-index ($1\\sigma$ upper-limits). The red dot-dashed line shows the $1\\sigma$ upper-limit of metallicity ($N2$-index) for the composite spectrum in $K$ band. We additionally plot the average LAE (pentagon: \\citealt{nakajima2012}) and two $z\\sim 2.4$ LAEs (triangles: \\citealt{finkelstein2011}). LAEs at $z=0.195-0.44$ are also shown with gray enclosed area \\citep{cowie2011}. For non-LAEs, we plot lensed galaxies at $1.5<z<2.5$ (open diamonds: \\citealt{richard2011}), the metal-poor BX galaxy at $z=2.3$ (open pentagon: \\citealt{erb2010}), and BX/BM galaxies (open squares: \\citealt{erb2006a}). The M-Z relation compiled by \\citet{maiolino2008} at $z\\sim 0$ and $z\\sim 2$ are drawn as long-dashed and dashed curves, respectively. All data plotted here have been recalibrated to have the same metallicity scale \\citep{maiolino2008} and IMF \\citep{salpeter1955} except for our LAEs with $R23$-index whose metallicities are calibrated with \\citet{KD2002}'s relations. (right) Fundamental metallicity relation proposed by \\citet{mannucci2010}. The dashed curve and the shaded areas indicate the relation and its typical dispersions respectively, determined in the local universe \\citep{mannucci2010,mannucci2011}. The open squares show $z\\sim 2$ star forming galaxies collected/compiled by \\citet{mannucci2010} from \\citet{erb2006a} and so on. Other symbols are the same as those in the M-Z relation. \\label{fig:MZR_FMR}} \\end{figure*}   Although exact geometries of \\HII\\ regions in LAEs are sill uncertain, the LAE's optically thin property has been suggested by recent studies. E.g., \\citet{hashimoto2013} investigate the kinematics of the galaxy scale outflows for LAEs using both \\Lya\\ emission and strong metal absorption lines, finding a tendency that LAEs have smaller values of neutral hydrogen column density in contrast with LBGs. Theoretically, \\citet{yajima2012} also suggest a completely ionized, optically thin phase can exist for LAEs. Therefore, the high ionization parameters may suggest LAEs can possess higher \\fesc, in contrast with LBGs showing modest fractions of photon escape. This trend is roughly consistent with the findings by \\citet{iwata2009} that LAEs tend to show smaller observed ratios of non-ionizing UV to Lyman continuum flux density than LBGs, as well as with the latest measurements of \\fesc s for both populations \\citep{nestor2013}. Similar suggestions of high escape fractions are also offered for LCGs, likely local analogs of high-$z$ emission line galaxies \\citep{jaskot2013,shim2012}. Since the fraction of LAEs among star-forming galaxies is known to increase with redshift (e.g., \\citealt{ouchi2008, stark2010}), LAEs and similar low-mass galaxies could produce ionizing photons efficiently and play a key role in supplying ionizing photons for the cosmic reionization at $z\\gtrsim 6$.   We note in conclusion that if the photon escapes are non-zero, LAEs' ionizing fluxes need not necessarily to be as hard as we expect for zero escape fraction. A hard ionizing photon flux (young population), a high escape fraction of ionizing photons, or a combination of both conditions likely cause the high ionization parameters found in the LAEs.      \\subsubsection{Metallicity} \\label{sssec:MZR}   Figure \\ref{fig:MZR_FMR} (left) shows the observed mass-metallicity (M-Z) relation. The red symbols represent LAEs at $z\\sim 2$, while the blue symbols show the other galaxies such as LBGs at the similar redshifts. Although some of our LAEs have just weak upper-limits of metallicity due to the lack of \\NII-detection (\\S\\ref{ssec:qion_Z}), for CDFS-3865 and COSMOS-30679, we obtain the metallicity estimates from $R23$-index with the oxygen lines. We plot these $R23$ metallicity estimates with larger circles in Figure \\ref{fig:MZR_FMR}. Albeit with its relatively large error in metallicity, CDFS-3865 falls below the conventional M-Z relation of $z\\sim 2$ LBGs. At least one LAE (HPS194) found by \\citet{finkelstein2011} also looks less chemically-enriched for its mass. Although the offsets of these LAEs from the M-Z relation can be due to the intrinsic scatter of the relation seen in the local universe (e.g., \\citealt{tremonti2004}), it may also indicate that they are less chemically evolved for their stellar masses. The idea is consistent with the suggestion of CDFS-3865's high ionization parameters. Since galaxies with high ionization parameters (e.g., BX418) tend to fall below the M-Z relation as well% \\footnote{ Galaxies with high ionization parameter \\citep{fosbury2003, richard2011} are omitted due to their high redshifts ($z>2.5$). Since the M-Z relation is known to evolve with redshift, a direct comparison with the $z\\sim 2$ LAEs is not appropriate.}, young galaxies with high ionization parameters may not follow the M-Z relation defined by more evolved galaxies. Unfortunately, no constraint on ionization parameter is given for HPS194. Alternatively, differences in star-formation activity may cause the scatter. We discuss this point later.   By contrast, the comparison with COSMOS-30679 is more unclear. Its large errors both in stellar mass and metallicity prevent us from telling if it is below/on the LBGs' relation. Since its ionization parameter is comparable to those for LBGs (\\S\\ref{sssec:high_qion}), COSMOS-30679 may be an LBG-like galaxy.   We also constrain an average metallicity from the $K$ band composite spectrum (Figure \\ref{fig:spec_ALL_n6}) by using the empirical \\NII-index. The $1\\sigma$ ($2\\sigma$) upper-limit of metallicity is \\Oabundance\\ $<8.42$ ($8.66$), which corresponds to $Z<0.54$ $(0.93)\\,Z_{\\odot}$. On the other hand, we independently obtain an average lower-limit of metallicity for LAEs to be \\Oabundance\\ $>7.93$ ($7.63$), or $Z>0.17\\,(0.09)\\,Z_{\\odot}$ at the $1\\sigma$ ($2\\sigma$) level, based on the \\OII/(\\Ha+\\NII) ratio whose fluxes are obtained by stacking $1.18$ and $2.09\\,\\mu$m narrowband images for more than $100$ LAEs \\citep{nakajima2012}. LAEs thus typically have a metallicity \\Oabundance\\ $=7.93$--$8.42$ ($7.63$--$8.66$) at the $1\\sigma$ ($2\\sigma$) level. The range is robust in the sense that the upper-limit is constrained by bright, massive LAEs while the lower-limit by faint, low-mass LAEs. More spectroscopic data are needed, however, to conclude that there are no exceptional LAEs that have nearly zero or super-solar metallicities. The metallicity range also suggests that LAEs at $z\\sim 2$ are rather less chemically enriched than those at $z\\sim 0.3$ (\\Oabundance\\ $\\sim 8.4$% \\footnote{ We recalculate the metallicity by using the \\citet{maiolino2008} indicator. The original estimate is $\\sim 0.15$\\,dex lower. }; \\citealt{cowie2011}).   \\citet{nakajima2012} find that LAEs fall typically above the M-Z relation below the stellar mass $\\sim 10^9\\,M_{\\odot}$. In contrast, the current study finds less chemically enriched LAEs for their masses. This apparent inconsistency may be due to the sampling of a large variety of evolutionary phases within the LAE population. Such variation has been indeed reported by other studies (e.g., \\citealt{nilsson2011,oteo2012}). Alternatively, differences in star-formation activity may cause the inconsistency. To check the possibility, we plot the LAEs on the fundamental metallicity relation (FMR; \\citealt{mannucci2010}), the relation between stellar mass, metallicity, and SFR. In Figure \\ref{fig:MZR_FMR} (right), most of the galaxies at $z=0$--$2$ including the average LAE \\citep{nakajima2012} appear to be consistent with the same FMR determined by the SDSS galaxies within their errors. However, some LAEs such as HPS194, and possibly CDFS-3865 and COSMOS-30679, still appear to fall below the relation. The inconsistency for HPS194 is remarkable ($\\gtrsim 5\\sigma$ level). Their low metallicities are not just likely due to their relatively high SFRs. We can speculate that the FMR is not universal and may fail to reproduce the properties of galaxies with e.g., high ionization parameter. Clearly however, much more data is needed to test the idea statistically.      \\subsubsection{Star-Formation Activity} \\label{sssec:MsSFRR}  \\begin{figure} \\epsscale{1.15} \\plotone{Mass_sSFR.eps} \\caption{ Relation between stellar mass and specific star formation rate (sSFR) for $z\\sim 2$ LAEs (red) and continuum selected galaxies (blue); LAEs with NIR spectroscopy (circles: this work, triangles: \\citealt{finkelstein2011}), the average LAE obtained by stacking NIR NB imaging (pentagon: \\citealt{nakajima2012}), BX/BM galaxies (open squares: \\citealt{erb2006b}), and sBzK galaxies (solid squares: \\citealt{hayashi2009}). COSMOS-43982, a possible object with AGN activity, is marked with black diamond. The red open circle denotes COSMOS-30679, which has less accurate SED fitting results due to the neighbor's contamination. All sSFRs plotted here are derived from \\Ha\\ luminosity with Salpeter IMF. The short horizontal bars in each symbols for the LAEs indicate the sSFR with no dust extinction correction. The dashed lines correspond to constant SFRs of $0.1, 1.0, 10, 100,$ and $1000\\,M_{\\odot}\\,{\\rm yr}^{-1}$. \\label{fig:MsSFRR}} \\end{figure}   In order to examine the star-formation activity of LAEs, we plot the LAEs on the sSFR versus stellar mass plane in Figure \\ref{fig:MsSFRR}. We also plot BX/BM galaxies \\citep{erb2006b} and sBzK galaxies \\citep{hayashi2009}. We note that this plot should be interpreted with cares since the spectroscopic data introduce limits in the sensitivity to low SFRs.   BX/BM and sBzK galaxies follow a simple scaling relation between sSFR and the stellar mass, whose tight relation is referred to as {\\it the star formation main sequence} (e.g., \\citealt{daddi2007}). Note that BX/BM and sBzK galaxies appear to have slightly different sequences; BX/BMs show lower sSFRs at a given mass. Compared to them, LAEs appear to follow the BX/BMs' main sequence almost over the full mass range. CDFS-3865 and COSMOS-30679, which exhibit high ionization parameters and low metallicities, are not outliers on this diagram. This trend indicates that star-formation activities are well determined by their stellar mass, irrespective of the presence of \\Lya\\ emission. Interestingly, \\citet{rhoads2013} recently find that LAEs tend to have higher SFR surface densities for their gas mass surface densities than ``normal'' star-forming galaxies including BzK galaxies.      \\subsection{Physical Properties inferred from \\Lya\\ and \\Ha\\ emission} \\label{ssec:Lya_Ha}  \\subsubsection{Correlation between EW(\\Lya) and EW(\\Ha)} \\label{sssec:EW_Lya_Ha}  Figure \\ref{fig:EW_Lya_Ha} shows equivalent widths of \\Lya\\ and \\Ha\\ for LAEs (hereafter referred to as the ``EWs diagram''). This plot may be useful to understand star-formation histories of LAEs, because the EWs' continuum fluxes evolve in different ways when different star-formation histories are assumed. Another advantage of using the EWs is that both are pure observables and to zeroth order independent of dust extinction (the effect of dust will be discussed at the end of this section).   The superposed curves on Figure \\ref{fig:EW_Lya_Ha} illustrate evolutions of the EWs for the two extreme star-formation histories, instantaneous burst (dashed) and constant star-formation (solid) at several metallicities \\citep{schaerer2003}% \\footnote{ The models are collected from the Strasbourg astronomical Data Center (CDS). We present here the three metallicity cases assuming Salpeter IMF \\citep{salpeter1955} with upper (lower) mass cut-off to be $100\\,M_{\\odot}$ ($1$\\,$M_{\\odot}$). Case B recombination is assumed for an electron temperature of $T_e=3\\times 10^4$\\,K at zero metallicity and $T_e=10^4$\\,K otherwise, and an electron density of $n_e=10^2$\\,cm$^{-3}$.}. For the instantaneous burst, since very massive stars ($M_{\\star}\\gtrsim 10\\,M_{\\odot}$) complete their evolution within $\\lesssim 10$\\,Myr, both EWs decline rapidly. As a result, their curves evolve quickly to lower-left on the EWs diagram. For the constant star-formation, on the other hand, the EW(\\Lya) stops declining around $\\sim 100$\\,Myr, because massive stars that are responsible for both \\Lya\\ emission and UV-continuum reach a steady mode. Since the EW(\\Ha) keeps declining as the older stars build up in the galaxy, their slopes of the evolutionary tracks become less steep on the EWs diagram. Although \\citeauthor{schaerer2003}'s calculations stop at $\\sim 400$\\,Myr for constant star-formation, the EW(\\Lya) varies little after $100$\\,Myr \\citep{CF1993}. Therefore, the tracks must extend to the left almost horizontally after the terminal points, reaching EW(\\Ha) $\\sim 100$\\,\\AA\\ for the solar-metallicity case when an age of $\\sim 3$\\,Gyr (e.g., \\citealt{leitherer1999})% \\footnote{ Although \\citet{leitherer1999} show the EW(\\Ha) evolution with age until $1$\\,Gyr, we extend the evolution toward older age smoothly to obtain EW(\\Ha) $\\sim 100$\\,\\AA\\ at $\\sim 3$\\,Gyr. We have confirmed the validity of the smooth extrapolation for constant star-formation by running GALAXEV \\citep{BC2003}. }.   Compared with the model tracks, some LAEs with EW(\\Ha) $\\sim 1000$\\,\\AA, CDFS-3865, HPS256 \\citep{finkelstein2011}, and possibly COSMOS-08501, can be explained by the instantaneous burst models with very young ages (a few Myr). The other LAEs whose EW(\\Ha) is modest appear to prefer the constant star-formation models with $\\gtrsim 100$\\,Myr.  Although the sample is small, we find more than half of the LAEs appear to need a continuous star-formation history rather than an instantaneous burst.   Note, however, that since the EWs are sensitive to recent starbursts, LAEs on instantaneous tracks are not necessarily very young, but can be experiencing a burst after a continuous star-formation. Such a combination of burst plus continuous star-formation is indeed needed to explain the EWs of \\Ha\\ and optical colors observed in local dwarf galaxies \\citep{lee2006}. An LAE with such a combined star-formation history will be on a track of instantaneous burst during starburst phases and on a track of constant star-formation for the remaining time.   \\begin{figure} \\epsscale{1.15} \\plotone{EW_Lya_Ha.eps} \\caption{ Relation between \\Lya\\ and \\Ha\\ equivalent widths. The red symbols show $z\\sim 2$ LAEs (same as in Figure \\ref{fig:MsSFRR}), and the gray squares show LAEs at $z=0.3$ \\citep{cowie2011}. Superposed lines are evolutions of the EWs for instantaneous burst (dashed) and constant star-formation (solid) at metallicities of solar (orange), sub-solar (green), and zero (blue) calculated by \\citet{schaerer2003}. Ages are denoted by the numbers (6, 7, 8) near the dots on the lines which indicate ($1$\\,Myr, $10$\\,Myr, $100$\\,Myr). For the solar-metallicity models, the lower point labeled \"6\" is for the instantaneous burst. \\label{fig:EW_Lya_Ha}} \\end{figure}  \\begin{figure*} \\epsscale{1.15} \\plotone{fesc.eps} \\caption{ (left) \\Lya/\\Ha\\ flux ratio vs. EW(\\Lya). Symbols are the same as in Figure \\ref{fig:EW_Lya_Ha}. The darker symbols show the ratios calculated by the observed \\Lya\\ luminosity divided by the intrinsic \\Ha\\ luminosity (i.e., lower-limit of the ratios), while the lighter show the ratios calculated by the observed \\Lya\\ and \\Ha\\ luminosities. The ratios for $z\\sim 0.3$ LAEs are based on observed luminosities. Dotted curve represents a single power-law fit given in Eq. (\\ref{eq:Flux_LyaHa_EW_Lya}). The dashed line and the gray shaded area show the \\Lya/\\Ha\\ flux ratio assuming the Case B recombination (8.7; \\citealt{brocklehurst1971}) and its variation with electron density ($8.1$--$9.2$ with $n_e=10^2$--$10^{10}$\\,cm$^{-3}$; \\citealt{HS1987}) when $T_e=10^4$\\,K, respectively. \\% (right) Escape fraction of \\Lya\\ photons (assuming the Case B) vs. dust extinction. The red squares show $z=2.2$ LAEs \\citep{hayes2010}, and the faint-colored red error bars show $z=2$--$4$ LAEs \\citep{blanc2011}. Superposed curves show the relations at a given $q$ parameter ($q=0.0, 0.1, 0.5, 1.0, 2.0, 10.0$). \\label{fig:fesc}} \\end{figure*}   Similarly, most of $z\\sim 0.3$ LAEs appear to prefer continuous star-formation on the EWs diagram. Interestingly, $z\\sim 2$ LAEs appear to have systematically higher EW(\\Lya) and EW(\\Ha) than lower-$z$ LAEs. Although our spectroscopic sample may be biased toward larger EW(\\Lya), the difference may be a sign that higher-$z$ LAEs are younger. This idea is supported by the inference of lower metallicities at higher-$z$ (\\S\\ref{sssec:MZR}; see also e.g., \\citealt{finkelstein2009,cowie2010,cowie2011}).   However, one worry on the use of the EWs diagram is the effect of dust on \\Lya. If the degree of dust extinction is different for \\Lya\\ and UV-continuum, EW(\\Lya) is no longer independent of dust extinction. In order to examine this effect, we introduce a $q$ parameter following \\citet{finkelstein2008}. The $q$ parameter is defined as $q=\\tau({\\rm Ly}\\alpha)/\\tau_{1216}$, where $\\tau({\\rm Ly}\\alpha)$ and $\\tau_{1216}$ are optical depth for \\Lya\\ and UV-continuum at $\\lambda=1216$\\,\\AA, respectively. Small $q$ values ($<1$) mean \\Lya\\ photons suffer less attenuation by dust than UV-continuum photons, while large values mean \\Lya\\ photons are more heavily attenuated. In the former (latter) case data points go down (up) on the EWs diagram after the corrections. From previous works, LAEs at $z\\sim 2$ have modest $q$ values; e.g., from the values given in \\citet{hayes2010} we calculate $q\\simeq 1$--$1.5$ for $z=2.2$ LAEs whose \\Lya\\ and \\Ha\\ luminosities are estimated from two narrowbands. Similarly, \\citet{nakajima2012} obtain $q=0.7\\pm 0.1$. \\citet{blanc2011} obtain $q=0.99$ for $z=2$--$4$ LAEs whose intrinsic \\Lya\\ luminosities are inferred from UV-continuum. Therefore, LAEs are on average likely to show $q\\sim 1$, and whose observed EWs(\\Lya) to be approximately intrinsic. However, we find a non-negligible scatter in $q$-values around unity (especially toward smaller values) for the individual LAEs with \\Ha\\ measurement (\\S\\ref{sssec:superCaseB}; Figure \\ref{fig:fesc} right). The LAEs with $q<1$ go down on the EWs diagram and approach the instantaneous burst tracks. Unfortunately, with the current large errors in $q$-values, we cannot clearly tell which star-formation histories are likely for LAEs. Future radiative transfer calculations as well as higher S/N ratio spectra of \\Lya\\ and \\Ha\\ (and probably \\Hb) will enable full use of the EWs diagram.      \\subsubsection{super Case B objects?} \\label{sssec:superCaseB}  Figure \\ref{fig:fesc} (left) shows the \\Lya/\\Ha\\ ratio against EW(\\Lya) for LAEs. The dark red symbols indicate the observed \\Lya\\ luminosity divided by the intrinsic \\Ha\\ luminosity corrected using the attenuation inferred from the SED fitting (i.e., showing lower-limits on the $y$-axis), while the light red symbols show the observed \\Lya/\\Ha\\ ratios. A trend that LAEs with larger EW(\\Lya) have larger \\Lya/\\Ha\\ ratio appears to be present. This trend itself is not so surprising, but interestingly some LAEs, COSMOS-08501, and possibly COSMOS-13636 and one LAE at $z\\sim 0.3$, may have \\Lya/\\Ha\\ ratios exceeding the Case B recombination value ($8.7$; \\citealt{brocklehurst1971}), which we call ``super Case B''. Super Case B \\Lya/\\Ha\\ ratios have also been reported for some strong LAEs in the local universe (\\citealt{hayes2007}, \\citealt{atek2008}, \\citealt{otifloranes2012}, and references therein). The variation of the intrinsic \\Lya/\\Ha\\ ratio (the gray shaded area) does not seem to be a significant issue. Although it is not obvious that super Case B objects are really included in our sample due to their large errors, it is worth discussing possible physical origins of them in case they really exist.   We consider the possible effect of geometry, and the kinematics of dust and gas in the ISM. \\citet{neufeld1991} propose a clumpy, multi-phase ISM where gas and dust are gathered in clouds within a low-density medium. With such circumstances, \\Lya\\ photons can be scattered at the surfaces of the clouds due to the resonant nature, while continuum photons would penetrate the clouds deeply. Since dust is contained in the clouds, \\Lya\\ photons would have a much smaller chance of encountering dust than other photons. In this scenario, a large EW(\\Lya) and \\Lya/\\Ha\\ ratio can be observed. Alternatively, an outflow of the ISM can be a cause of strong \\Lya\\ emission (e.g., \\citealt{kunth1998}). However, \\citet{hashimoto2013} find an anti-correlation between EW(\\Lya) and \\Lya\\ velocity offset for LAEs, the latter is considered to be positively correlated with outflow velocity. This anti-correlation is also supported by the observations of LBGs (\\citealt{adelberger2003}; see also \\citealt{shapley2003,pettini2001}). In addition, the authors measure the LAEs' outflow velocity directly from metal absorption lines, finding that outflow velocities for LAEs and LBGs are comparable. Therefore, outflows do not appear to be a major mechanism for producing large EW(\\Lya) and \\Lya/\\Ha\\ ratio, though the sample size of LAEs is still small and much more data as well as theoretical works to interpret the properties are needed.   In order to clarify the effects of a potentially clumpy geometry of ISM on \\Lya/\\Ha\\ ratios, we plot in Figure \\ref{fig:fesc} (right) the relation between $E(B-V)$ and the escape fraction of \\Lya\\ photons (\\fescLya) under the Case B recombination assumption. We estimate \\fescLya\\ as \\begin{eqnarray} f_{\\rm esc}^{{\\rm Ly}\\alpha} \\equiv \\frac{L_{\\rm obs}({\\rm Ly}\\alpha)}{L_{\\rm int}({\\rm Ly}\\alpha)} = \\frac{L_{\\rm obs}({\\rm Ly}\\alpha)} {8.7 L_{\\rm int}({\\rm H}\\alpha)}, \\label{eq:Lya_f_esc} \\end{eqnarray} where subscripts {\\lq}int{\\rq} and {\\lq}obs{\\rq} refer to the intrinsic and observed quantities, respectively. The intrinsic \\Ha\\ luminosities are derived from the observed \\Ha\\ fluxes, corrected for dust extinctions. The superposed lines show the relations at a given $q$ parameter (\\S\\ref{sssec:EW_Lya_Ha}); \\begin{eqnarray} q = \\frac{-\\log \\left(f^{{\\rm Ly}\\alpha}_{\\rm esc}\\right)} {0.4k_{1216}E(B-V)}, \\label{eq:q} \\end{eqnarray} where $k_{1216}$ is an extinction coefficient at $\\lambda=1216$\\,\\AA\\ (11.98; \\citealt{calzetti2000}). The clumpy geometry of ISM (or outflow) is favored by objects with $q=0$--$1$. From Figure \\ref{fig:fesc} (right), most of the LAEs presented here are located in the range $q=0$--$1$; e.g., COSMOS-13636 and COSMOS-30679 have large \\fescLya\\ in spite of their moderate amounts of dust. SSA22-8043 has a large $q$ parameter of $\\sim 10$, and can be an exception if its \\Lya\\ is heavily resonant-scatted by neutral hydrogen gas.   A notable object is COSMOS-08501. It appears to fall above the $q=0$ line, where the clumpy ISM model does not work assuming Case B recombination, although the errors are relatively large. Moreover, since it is inferred to possess relatively small amount of dust, the large EW(\\Lya) owing to the clumpy ISM is unlikely. In case the object is really super Case B, the only remaining explanation is the \\Lya\\ enhancement caused by collisional excitations. Due to the decreasing collisional strengths with increasing principle quantum number, collisional excitations can lead to \\Lya/\\Ha\\ ratios over the Case B value (see also \\citealt{osterbrock1989}). Shocks caused by interactions with other sources, AGN activity, supernova explosions, strong outflows or infall are possible candidates for the collisional excitations.   Based on the HST images (Figure \\ref{fig:acs_I}), COSMOS-08501, a super Case B candidate, looks very compact and shows no sign of interactions. COSMOS-13636, which has a very small $q$ parameter, shows two faint sources nearby within $\\sim 5$\\,kpc (projected) from the object. Its strong \\Lya\\ emission can be (partly) due to shocks caused by interactions. Indeed, some fractions of LAEs are turned out to exhibit morphologies suggestive of mergers at higher redshift (e.g., \\citealt{pirzkal2007,bond2009}) as well as lower redshift \\citep{cowie2010}. Theoretically, \\citet{tilvi2011} demonstrate that mergers play an important role in mass assembly and star-formation in majority of the LAEs especially at higher redshift. Unfortunately, the shock-induced \\Lya\\ emission due to mergings has not yet been taken into account in these theoretical works (see also \\citealt{tilvi2009}). Alternatively, as discussed by \\citet{mori2004}, supernova explosions can cause strong shocks, resulting in strong \\Lya\\ emission. Although the authors intend to explain extended \\Lya\\ blobs ($\\sim 100$\\,kpc) with high \\Lya\\ luminosities ($\\sim 10^{43}$\\ergs), their basic ideas can be applied to normal LAEs. AGN activity and outflows seem unlikely (see \\S\\ref{ssec:AGN} and \\citealt{hashimoto2013}).   When shocks (even partly) contribute to emission lines of LAEs, estimates of physical quantities such as ionization parameter (from \\OIII/\\OII\\ ratio), metallicity (from $N2$-index), SFR (from \\Ha), and dust extinction (from Balmer decrement) become less accurate. In particular, the intrinsic \\Ha/\\Hb\\ value becomes larger when shocks are present, resulting in an overestimate of the abundance of dust. This effect may also help explain the presence of super Case B objects (e.g., \\citealt{otifloranes2012}). Thus, the results presented in this paper may require some corrections for the presence of shocks. We plan to address this in future work, through the simultaneous application of photo-ionization and shock models, with deeper spectroscopy which will detect the weaker lines not fully detected in the observations presented here (e.g., Balmer lines, \\OII, \\OIII, \\NII, \\SII).   A final remark is that from our discussions so far it is evident that \\Lya\\ is not a robust indicator of SFR for LAEs. The observed data ($z=0$--$2$) in Figure \\ref{fig:fesc} (left) are relatively well represented by a single power-law fit \\begin{eqnarray} \\log \\left({\\rm Ly}\\alpha/{\\rm H}\\alpha\\right)_{\\rm obs} = && (-1.08\\pm 0.29) \\nonumber \\\\ && + (0.87\\pm 0.19)\\times \\log {\\rm EW}({\\rm Ly}\\alpha), \\label{eq:Flux_LyaHa_EW_Lya} \\end{eqnarray} which is shown by the dotted curve in Figure \\ref{fig:fesc} (left). According to this simple relation, SFRs of LAEs with EW(\\Lya) $\\sim 20$\\,\\AA\\ (a typical threshold in narrowband searches) can be underestimated by a factor of about $10$. Thus, SFRs estimated from \\Lya\\ may involve a factor of $\\sim 10$ uncertainties intrinsically.        We have presented NIRSPEC and MMIRS rest-frame optical spectra of seven \\Lya\\ emitters (LAEs) at $z=2.2$, which are selected from our Subaru/Suprime-Cam NB387 survey in COSMOS, Chandra Deep Field South, and SSA22. Our first NIR spectroscopic result discusses the kinematics of LAEs and is presented in \\citet{hashimoto2013}. As a companion study, this paper presents mainly the ionization and chemical properties of LAEs based on multiple nebular lines. Our sample includes one possibly AGN-dominated galaxy, and six star-forming galaxies. \\Ha\\ is detected in all six star-forming LAEs, while \\NII$\\lambda6584$ is only detected in the galaxy with signs of AGN activity. Among the six star-forming galaxies, one (CDFS-3865) also has detections of \\OII$\\lambda3727$, \\Hb, and \\OIII$\\lambda\\lambda 5007, 4959$, and another (COSMOS-30679) has detections of \\OII\\ and \\OIII. Our deep $J$ band spectroscopic observations provide the first \\OII-detections for two individual LAEs at high-$z$. Our main results are summarized as follows.   \\begin{itemize}  \\item % The \\OIII/\\OII\\ ratio vs. $R23$-index diagram reveals that CDFS-3865 has a very high ionization parameter ($q_{ion}=2.5^{+1.7}_{-0.8} \\times 10^8$\\,cm\\,s$^{-1}$) and a low oxygen abundance (metallicity; \\Oabundance\\ $=7.84^{+0.24}_{-0.25}$) in contrast with moderate values of other high-$z$ galaxies such as LBGs. COSMOS-30679 also has a relatively high ionization parameter ($q_{ion}=8^{+10}_{-4}\\times 10^7$\\,cm\\,s$^{-1}$) and a low metallicity (\\Oabundance\\ $=8.18^{+0.28}_{-0.28}$). LAEs would therefore 1) represent an early stage of galaxy formation dominated by massive stars in compact star-forming regions, and/or 2) have a higher escape fraction of ionizing photons probably achieved by density-bounded \\HII\\ regions. High-$q_{ion}$ galaxies like LAEs would thus play a key role in supplying ionizing photons for cosmic reionization in the early universe.   \\item % Local low-metallicity galaxies ($7\\lesssim$ \\Oabundance\\ $\\lesssim 8.5$) show similar ionization parameters and metallicities to high-$z$ star-forming galaxies, and those with extreme quantities are likely analogs of high-$z$ LAEs. One such population is ``green pea'' galaxies (GPs; \\citealt{cardamone2009}), in terms of their low-metallicity, high \\OIII/\\OII\\ ratios, and relatively high SFRs. A notable difference between the GPs and LAEs are their abundances; the GPs occupy only $\\sim 0.06$\\,\\% of the SDSS galaxy sample, and its number density is almost two order of magnitude smaller than that of LAEs at $z\\sim 2$.   \\item % CDFS-3865 falls below the mass-metallicity relation of LBGs at similar redshifts. Its low metallicity seems not to be explained by its star formation rate being taken into account, albeit with its relatively large error. COSMOS-30679 appears to exhibit the same trend, although its large error in metallicity complicates the interpretation. Interestingly, galaxies with high ionization parameters tend to fall below the relation. Such galaxies may not follow the relation determined by more evolved galaxies.   \\item % The composite spectrum independently provides an upper-limit on the metallicity of \\Oabundance\\ $<8.42$ ($<8.66$) at the $1\\sigma$ ($2\\sigma$) level. Combined with an lower-limit of metallicity \\citep{nakajima2012}, LAEs typically have metallicities \\Oabundance\\ $=7.93$--$8.42$ ($7.63$--$8.66$) at the $1\\sigma$ ($2\\sigma$) level.   \\item % In contrast to the large differences in ionization parameters and metallicity between LAEs and LBGs, we find LAEs have similar specific star formation rates as BX/BM galaxies at a given stellar mass.   \\item % The EW(\\Lya) vs. EW(\\Ha) diagram interestingly suggests that more than half of the LAEs appear to need an extended star-formation such as a burst superimposed upon a continuous star-formation rather than the instantaneous burst alone. However, since EW(\\Lya) may suffer from effects of dust and we do find a non-negligible scatter in $q$-values around unity (especially toward smaller values), we need to carefully interpret the result.   \\item % LAEs with low $q$-values ($q=0$--$1$) can be explained by the clumpy geometry of ISM. Interestingly, our sample may include objects with further enhanced \\Lya, which we call super Case B. If they really exist, the only possible explanation is the collisional excitations of \\Lya. Interactions with other sources and/or supernova explosions are possible key events that may cause shock-induced collisional excitation. If such shocks play a role in enhancement of the \\Lya\\ flux, physical quantities such as ionization parameter, metallicity, SFR, and dust extinction should be re-computed using a combination of photo-ionization and shock-excitation models. We plan to investigate the role of shocks further in future works.  \\end{itemize}"
    },
    "1208/1208.1579_arXiv.txt": {
        "abstract": "We have investigated the color-magnitude diagram of $\\omega$ Centauri and find that the blue main sequence (bMS) can be reproduced only by models that have a of helium abundance in the range $Y=0.35$--$0.40$.  To explain the faint subgiant branch of the reddest stars (``MS-a/RG-a\" sequence), isochrones for the observed metallicity ([Fe/H] $\\approx -0.7$) appear to require both a high age ($\\sim 13$ Gyr) and enhanced CNO abundances ([CNO/Fe] $\\approx 0.9$). $Y \\approx 0.35$ must also be assumed in order to counteract the effects of high CNO on turnoff colors, and thereby to obtain a good fit to the relatively blue turnoff of this stellar population. This suggest a short chemical evolution period of time ($< 1$ Gyr) for \\ocen. Our intermediate-mass (super-)AGB models are able to reproduce the high helium abundances, along with [N/Fe] $\\sim 2$ and substantial O depletions if uncertainties in the treatment of convection are fully taken into account. These abundance features distinguish the bMS stars from the dominant [Fe/H] $\\approx -1.7$ population.  The most massive super-AGB stellar models ($\\mzams\\geq6.8\\msun$, $M_\\mathrm{He,core} \\geq 1.245\\msun$) predict too large N-enhancements, which limits their role in contributing to the extreme populations. In order to address the observed central concentration of stars with He-rich abundance we show here quantitatively that highly He- and N-enriched AGB ejecta have particularly efficient cooling properties.  Based on these results and on the reconstruction of the orbit of \\ocen\\ with respect to the Milky Way we propose the galactic plane passage gas purging scenario for the chemical evolution of this cluster. The bMS population formed shortly after the purging of most of the cluster gas as a result of the passage of \\ocen\\ through the Galactic disk (which occurs today every $\\sim 40\\mem{Myrs}$ for \\ocen) when the initial-mass function of the dominant population had ``burned\" through most of the Type II supernova mass range.  AGB stars would eject most of their masses into the gas-depleted cluster through low-velocity winds that sink to the cluster core due to their favorable cooling properties and form the bMS population. In our discussion we follow our model through four passage events, which could explain not only some key properties of the bMS, but also of the MS-a/RGB-a and the $s$-enriched stars. ",
        "introduction": "\\label{sec:intro} Omega Centauri provides an especially valuable constraint on our understanding of chemical evolution because of its unusual properties and its proximity, which enables us to study its component stellar populations in great detail over a very extended range in luminosity. While it has been known for some time that the giants in this system encompass a range in [Fe/H] from $\\sim -2.2$ to $-0.5$ \\citep[e.g.][]{brown:93,suntzeff:96}, the extensive spectroscopic surveys carried out in the past decade, in particular \\citep[][and references therein]{smith:00,hilker:04,kayser:06,vanloon:07,johnson:08,johnson:10,marino:12}, have established that the metallicity distribution rises sharply from [Fe/H] $\\approx -2.2$ to a strong peak at [Fe/H] $\\approx -1.7$, with a long tail that drops off to higher metal abundances containing secondary peaks at [Fe/H] $\\approx -1.4$, $-1.1$, and $-0.7$.  Type Ia supernovae appear to have contributed to the chemical makeup of only the most metal-rich stars given that their measured [$\\alpha$/Fe] abundance ratios are significantly reduced from the constant value of $\\approx 0.35$ found in stars having [Fe/H] $< -0.8$ \\citep{pancino:02,origlia:03}.  The elevated $\\alpha$-element abundances over most of the range in metallicity (also see Kayser et al.) implies that Type II supernovae were the major producers of these elements.  Moreover, the rapid rise in the [La/Fe] ratio between the most metal-deficient and the [Fe/H] $> -1.5$ populations (by a factor of 3 according to Johnson \\& Pilachowski) indicates that $s$-processing, and therefore, intermediate-mass asymptotic-giant branch (AGB) stars contributed significantly to the chemistry of $\\omega$ Cen after the dominant [Fe/H] $= -1.7$ population had formed.  Accompanying these spectroscopic advances have been equally impressive improvements in the photometric data.  The detailed, and deep, color-magnitude diagrams (CMDs) derived by \\citet{lee:99,hughes:00,rey:04,bedin:04,sollima:05a,sollima:07,villanova:07,calamida:09}, and \\citet{bellini:09,bellini:10}, among others, have also established the existence of several discrete stellar populations in $\\omega$ Cen.  The most baffling one of them is a so-called ``blue main sequence\" (bMS), discovered by \\citep{anderson:97}, that is clearly separated from, and bluer at a given magnitude than, the MS associated with the dominant [Fe/H] $\\approx -1.7$ population (see the CMDs reported by, e.g., Bedin et al.~and Villanova et al.).   What has made this discovery so puzzling is the determination by \\citet{piotto:05} that the former is more metal rich than the latter by $\\approx 0.3$ dex.  As discussed by Piotto et al., and anticipated by Bedin et al.~ and \\citet{norris:04}, the most obvious (only?) way to reconcile these observations with stellar evolutionary theory is to infer that the bMS stars have unusually high helium abundances ($Y > 0.35$).  Although it has not yet been possible to make a definitive connection of the bMS through the subgiant region of the CMD to the RGB \\citep[but see][]{king:12}, \\citet{johnson:10} have found that the giants in their sample with ``extreme\" abundances (i.e., those with low C and O, together with high N, Na, and Al; see their Fig.~23) constitute a similar proportion of all giants ($\\approx 27\\%$) as the fraction of all MS stars that are located on the bMS.  Furthermore, since (for the most part) their radial distributions are quite similar, it is tempting to conclude that the aforementioned giants are the descendants of stars that had occupied the bMS earlier in their evolutionary history.  However, because the O-poor stars do not exhibit a correlation with [La/Fe], while the O-rich stars do, and because there is no apparent correlation of the La abundance with radius, in contrast to the observed trend for the O-poor stars, Johnson \\& Pilachowski suggest that the observed light element abundances reflect both (i) the retention by $\\omega$ Cen of the ejecta of AGB stars and (ii) {\\it in situ} mixing on the RGB.  There is a second feature of the CMD that is potentially quite challenging to explain, and that is the extension of the reddest giant branch (independently discovered by \\citet{lee:99} and by \\citet{pancino:00}, who gave it the name ``RGB-a\" that has since been used to refer to it), through the subgiant branch \\citep{ferraro:04} to very faint magnitudes on the MS \\citep{bellini:10}.  Because this is is the most metal-rich of the discrete populations that have been identified, it is presumably also the youngest one.  Consequently, its age provides a key constraint on the timescale over which most of the chemical evolution took place in $\\omega$ Cen.  Because the turnoff of this fiducial sequence is so faint, the age of this stellar population must be quite old, implying that the bulk of the star formation occurred over a rather short interval of time.  Indeed, the reason why some investigations \\citep[e.g.][]{hilker:04,rey:04} derived an extended star formation history ($\\sim 3$--4 Gyr) in their analyses is that they had mistakenly adopted a bright turnoff for the RGB-a population.  However, it remains unclear whether the different stellar populations of \\ocen\\ have the same age to within a few $\\times 10^8$ yr \\citep{D'Antona2011,valcarce:11}, or they span a range in age of $\\lta 2$ Gyr \\citep{sollima:05a} or as much as 3--5 Gyr \\citep{stanford:06,villanova:07}.  Among stellar evolution models of different mass ranges, intermediate-mass and super-AGB stars have been proposed as the source of the He-enriched and otherwise extreme abundance patterns associated with the bMS. However, models of these types of stars do not always reproduce the required high He, high N, and low O abundance associated with the bMS in a quantitative way. For example, the intermediate-mass AGB models by \\citet[][specifically the $Z=\\natlog{6}{-4}$ case that is applicable to the first, most metal-deficient generation in \\ocen]{ventura:09} predict $Y=0.329$ and [N/Fe]$=1.84$ for $\\mzams=5.0\\,\\msun$ and $Y=0.360$ and [N/Fe]$=1.36$ for $\\mzams=6.0\\,\\msun$. The super-AGB models by \\citet{ventura:11} produce a maximum ejecta He abundance of $Y=0.36$, but information for N is not available.  The envelope He abundance at the end of the second dredge-up in the super-AGB models by \\citet{siess:07} reach $0.37$ for non-overshooting models, and $0.38$ for models with core overshooting. \\citet{siess:10} do not provide a model for $Z=\\natlog{6}{-4}$, but interpolating between their super-AGB thermal pulse calculations for $Z=0.001$ and $Z=10^{-4}$ yield He mass fractions in the ejecta of $Y\\sim 0.33$, while N is enhanced by a factor of $10$ to $100$. Moreover, those models with the largest increases in the nitrogen abundance also have the smallest O depletion, and models with O depletions of $\\sim 1 \\dex$ have N enhancements of only $\\sim 1 \\dex$.  While existing models clearly point qualitatively in the right direction, it is not clear to us what it would take to reproduce the ``extreme'' abundance mixtures that are associated with the bMS in \\ocen\\ in a quantitative sense.  Producing He-abundances in intermediate-mass AGB ejecta that exceeds the He-abundance of the ``extreme'' abundance mix in any significant way appears to be very difficult within model calculations. This does not leave a lot of room (if any) for dilution of the AGB ejecta before the formation of the blue main-sequence stars. Dilution has been suggested, however, in order to reproduce the abundance anti-correlations \\citep{d'ercole:08}. We are considering here rather the origin and evolution of one sub-population at a time. The presentation of Na-O abundances by sub-populations identified with a four-criterion cluster analysis by \\citet[][Fig.\\,6]{Gratton:2011kr} rather suggests that the anti-correlation is the result of the superposition of discrete Na and O abundance in different sub-populations, and as we will argue later (\\kap{sec:final-ONa}) dilution may not be needed. We therefore investigate the evolution of first-generation intermediate-mass AGB and super-AGB stars with the goal to generate the ``extreme'' abundance mix without taking into account any dilution.   In any case, the AGB ejecta may preferentially converge in the cluster center via AGB cooling flows \\citep[][and refs.\\ there]{DAntona:2011dm}. Such gas may have to be isolated, for example through supernova-induced clearing of intracluster gas that is ejected from stars having initial masses outside the range that encompasses intermediate-mass AGB and super-AGB stars \\citep{Charlie2011a}, while the tidal stripping of old stars increases the ratio of second- to first-generation stars \\citep{bekki:11}. However, complete SN gas purging makes it difficult to envisage how a spread in [Fe/H], like that observed in \\ocen, is produced.   The present study has been undertaken to address some of the issues described above.  In \\S 2, newly computed sets of isochrones and zero-age horizontal branch (ZAHB) loci for different values of $Y$, $Z$, and heavy-element mixtures are compared with the CMD of $\\omega$ Cen in order to illustrate how the interpretation of the observations would be affected by the assumed variations in the chemical abundances, and to make some assessment of the age and helium content of the stars that belong to the bMS and MS-a components.  \\S 3 investigates under which assumptions the predicted yields from $> 5\\msun$ AGB stars can be made to agree with the observed/inferred high $Y$, high-N, low-O abundances. Finally, a short summary of the main results of this study is given in \\S 4, which also discusses in some detail how our stellar evolutionary results together with the efficient cooling properties of the ejecta from intermediate-mass AGB stars may naturally explain the formation of additional populations like the bMS.  A key point in our proposed scenario is the periodic purging of the gas from globular clusters (or dwarf galaxies) that is expected to occur throughout their evolutionary histories whenever their orbits cause them to pass through the Galactic plane. Implications of this scenario are briefly summarized in \\S 5. ",
        "conclusions": "\\label{sec:conclusions}  \\subsection{Implications of the Galactic plane passage gas purging model} As \\citet{d'ercole:08} put it, He-rich populations (like the bMS in \\ocen) are only the tip of the iceberg of the phenomenon of second (or third, etc.) stellar populations. While our scenario of periodic gas purging events caused by the passages of clusters through the Galactic plane appears to be able to explain quite naturally the origin of the helium-rich bMS in \\ocen\\ (and possibly the more Fe-rich MS-a/RGB-a population, as well), one can readily imagine that this simple picture, if applied to clusters having a wide range of progenitor histories (e.g., masses and orbits), may, in fact, generate quite a variety of realizations in the globular clusters that we observe today.  Since each population would be associated with the ejecta arising from stars within one, or very few, relatively narrow initial mass ranges, isolated by successive Galactic plane passages that clear out any gas that had accumulated since the last passage, any two (or more) populations that are present in different clusters should not be identical, though possibly similar. The contributing segments of the IMF are statistically selected according to the progenitor orbit and the merger history.  However, in our scenario, Galactic plane passage intervals of $40$ to $100$ million years do favor the formation of a second, helium-rich population. If the delay is $40$ million years \\citep[as also found by][]{marino:12}, upper and lower mass limits can both be imposed by the Galactic plane passages. For time intervals as large as $100$ million years, only the lower mass cut-off for the progenitor stars of the second generation is set by the Galactic plane passage, while the higher mass cut-off would be due to SN purging. In fact, this case is similar to the model suggested by \\citet[][Sect.\\,3.1]{d'ercole:08}, who adopt $100$ Myr for their parameter $\\Delta t_\\mem{f}$. They do not mention what sets this time interval, aside from noting that, if $\\Delta t_\\mem{f}$ was longer, SN Ia and C-producing AGB stars would contribute to the second-generation star formation (see the above discussion). This is indeed why $\\Delta t_\\mem{f}$ should not be longer if the desired outcome is a helium-rich second generation, but we suggest that the main reason is that the Galactic plane passages which clear out the gas and terminate star formation occur every $\\lta 100$ Myr.  In summary, the simple principle of periodic Galactic plane passage purging in combination with low-velocity winds from massive AGB stars AGB wind predictions and their preferential cooling properties, may be able to account for the bMS abundances including its central concentration, the MS-a/RGB-a abundance patterns as well as the homogeneous distribution of and wide range of [Fe/H] observed in s-enhanced stars.  \\subsection{O-Na anti-correlation, n-capture element abundances, and other observed properties of \\ocen} \\label{sec:final-ONa} One of the \\citep[possibly defining,][]{Carretta:2010il} features of globular clusters is the more or less complete presence of the O-Na anticorrelation. How does it fit in with the Galactic plane passage gas purging model? The key question is whether (a) the anti-correlation occurs within a sub-population, or (b) it is the superposition of the rather distinct O-Na abundances of the present sub-populations. The answer will depend on how the latter are identified. If one uses only the [Fe/H] abundance \\citep[e.g.][]{DAntona:2011dm} one may indeed combine N-rich/C-poor with C-rich/N-poor stars into one group \\citep[][Fig.\\,2]{marino:12} and conclude that abundance anti-correlations are present even within a sub-population.  This view point may motivate, or even require as an explanationof the Na-O anti-correlation, a dilution scenario \\citep{DErcole:2011jca} which assumes that, while second-generation stars form in the cluster, unprocessed pristine gas is accreted from outside of the cluster. However, if option (b) is in fact the case, for example, if the sub-population identification purely by [Fe/H] is not entirely accurate, then an alternative interpretation of the anti-correlation is possible in which dilution may not play an important role.  Option (b) may also be favoured by observations of rather uniform Al abundances in the SGB-a population (in our scenario the third generation, see \\tab{tab:gppgpm}) by \\citet{pancino:11}, suggesting that no anti-correlation is present in this sub-population. Although we have tried to resist the temptation in this paper to apply the Galactic plane passage gas purging model to other clusters we note that \\citet{Carretta:2012jp} found that the anti-correlations involving Al, Mg, Na and O in NGC\\,6752 -- another GC for which multiple populations possibly including large He-enrichements have been found -- manifest themselves in a rather discrete fashion, in which different levels of enhancement and depletion cluster around discrete values that can be associated with individual sub-populations, rather than a continuous distribution that would be sugggested by a pure dilution mechanism.  The Galactic plane passage gas purging scenario implies that two distinct populations may have the same [Fe/H] abundance, and therefore that a collection of stars with the same [Fe/H] may not necessarily belong to the same population. As a matter of fact, many possible elemental markers if used by themselves may lead to degenerate grouping of stars. In order to identify those stars that most likely represent a coevally formed sub-population, several elemental markers should be combined. Such a ``group'' analysis, taking into account four abundance features simultaneously, has been performed by \\citet{Gratton:2011kr}. The resultant sub-population identification is indeed rather suggestive of option (b) mentioned in the previous paragraph (superposition of the distinct O-Na abundance markers from sub-population). Of course, there is always a concern that the particular choice of group-finding criteria biases the process. For that reason the sub-population identification by \\citet{Gratton:2011kr} may be evolve in the future when more observational data are added to this kind of analysis. In any case, this approach seems to be an improvement over just using [Fe/H] to identify sub-populations. As a consequence, we have no compelling need for the dilution mechanism in our model for \\ocen.  Our scenario has a few more implications -- some of which are summarized in \\tab{tab:gppgpm} -- that we would like to briefly discuss. It offers an alternative interpretation of the star formation time-scale determination of \\citet{DAntona:2011dm} that was based on the assumption that Fe in the MS-a/RGB-a populations originates in SN Ia. In our scenario, Fe for the \\emph{third} generation (MS-a/RGB-a) comes from the lowest mass, \\emph{second} generation SN II  after the second Galactic plane passage (\\kap{sec:scenario}), which would be in agreement with the $\\alpha$-element abundance patterns found by \\citet{Gratton:2011kr} for the most Fe-rich population labeled \\#2b. The time-scale limit from SN Ia would therefore not apply. In addition, star formation can take place even after the formation of this third generation from the late stellar winds of low-mass first-generation stars.  In fact, if we retain our chemical evolution time range determination from CMD considerations ($\\lta 1$ Gyr, \\kap{sec:cmd}), stars in the fourth generation can form from the slow wind ejecta of first-generation stars down to $1.8\\msun$. For stars with $\\apleq 2.6\\msun$ these ejecta would be C-rich (as well as La-rich from the main \\sprn). Because the first-generation donor stars are not centrally concentrated, the fourth generation stars that form out of them would also not be centrally located --- if the cooling efficiency of lower-mass AGB ejecta are smaller, or because the mass of the cluster is already smaller at this point. Indeed, \\citet{johnson:10} report no radial gradient for [La/Fe]. In any case, the observations reported by \\citet{marino:12} show that there are two different La-enrichment sequences, one associated with increasing C (O-rich/Na-poor group) and one associated with low O and high Na. In our scenario the former represents the transition from first to \\emph{fourth} generation, while the latter is associated with both the \\emph{second} and, possibly for even more extreme levels of enrichment, the \\emph{third} generation.  If we interprete the correlation of C with La in the O-rich/Na-poor group of \\citet{marino:12} as the result of the third dredge-up in $1.8$ -- $2.6\\msun$ \\emph{first} (and maybe \\emph{second} slightly higher [Fe/H]) generation stars, then their Fig.\\,7 implies that O and Na could also be the result of the third dredge-up. In fact, the slight increase of O with [Fe/H] in the O-rich/Na-poor group has been perceived to be quite a puzzle that implies an extra source of O. However, the models of \\citet[][Fig.\\,6]{herwig:04a} for $\\mathrm{[Fe/H] } = -2.3$ do indeed predict $\\mathrm{[C/Fe]} = 3.0$, $\\mathrm{[O/Fe]} = 1.8$ and $\\mathrm{[Na/Fe]} = 1.3$ for the average abundance in ejecta of a $2\\msun$ model. The O comes from the dredge-up of primary He-burning products that becomes appreciable at these low metal abundances, and such O enhancements are indeed observed in many CEMP stars that may carry the mass transfer signature from genuine low-metallicity AGB stars \\citep[for example,][]{sivarani:06,Kennedy:2011db}. The Na in the $2\\msun$ model comes from \\sprn\\ in the He-shell flash convection zone, where the neutron from the $\\nezw(\\alpha,\\n)\\mgfu$ reaction is captured again by \\nezw. The predicted enrichment levels suggest that a population which forms out of such ejecta should have about twice the enrichment of C compared to O and Na, which seems to be consistent with the \\citet{marino:12} data.  Star Leiden\\, 44462 has been tentatively identified by \\citet{DOrazi:2011jf} as a mass transfer object in order to account for the extremely high C abundance. Although radial velocity measurements may support this possibility, an alternative interpretation is that this star is part of the fourth generation (\\tab{tab:gppgpm}) which forms out of the slow winds of $1.8$ -- $2.6\\msun$ first (or second for higher [Fe/H]) generation stars. In fact, this star coincides very well with the C-La correlation sequence shown in \\citet[][Fig.\\, 7, upper right panel]{marino:12}.  But the O-poor/Na-rich group of \\citet{marino:12}, as well as all but the lowest [Fe/H] stars in \\citet{DOrazi:2011jf}, also show marked n-capture enhancements. Since in these stars the C abundance is low, the heavy elements cannot come from these lower-mass AGB stars. Instead we have to consider higher-mass AGB stars as well as super-AGB stars. The n-capture element predictions of \\citet{Karakas:2012kc}, e.g.\\ their  $6\\msun$, $Z=0.0001$ stellar evolution model, are based mostly on the \\nezw\\ neutron source in the He-shell flash convection zone (although some contribution from a \\cdr-pocket may be present as well). Clearly, the \\spr\\ models for these low-metallicity intermediate-mass stars are quantiatively still rather uncertain (cf.\\ \\kap{sec:increased_amlt}). However, the models of \\citet{Karakas:2012kc} do predict that higher mass AGB stars at this metal content do produce \\spr\\ elements, possibly with significant enrichment factors, and with a ratio of light (ls: Sr, Y, Zr, Rb) to heavy (hs: Ba, La) \\spr-elements that is higher than in the ejecta of lower mass stars. Such a signature would qualitatively agree with the n-capture abundances reported for the N-rich, intermediate- (bMS, second generation) and high- (MS-a/RGB-a, third generation) Fe group by \\citep{DOrazi:2011jf}.  In addition to the n-capture production by $\\nezw(\\alpha,\\n)\\mgfu$ in the He-intershell in intermediate mass AGB stars, we mentioned in \\kap{sec:inimass} the possible presence of \\ipr-conditions in super-AGB models with CBM \\citep{herwig:12b}, which would provide for another ``non-standard'' source of n-capture elements in stars that produce He-rich and O-poor populations in \\ocen. Both of these sources would be responsible for the La-enhanced sequence of the O-poor/Na-rich group reported by \\citet{marino:12}.  An important consequence of this discussion is the notion that the Na enhancements which we expect from AGB stars from $1.8\\msun$ all the way up to the super-AGB stars at $\\apleq 6.8$ -- $7.7\\msun$ may always be expected to go along with n-capture element enhancements. This assessment is supported by the La-Na correlations of both the O-rich/Na-poor and the O-poor/Na-rich group of \\citet[][lower-right panel of Fig.\\,7]{marino:12}.  \\subsection{Summary} \\label{sec:final_sum} Obviously the observational identification of the four generations that we have specified -- within the Galactic plane passage gas purging model -- as the result of highly idealized processes, is complicated by interference and superposition effects as well as contributions from, e.g., supernova purging, turbulent and tidal mixing, and cooling and mass loss flows, all of which are expected in a real cluster. In addition, observational uncertainties may cause some migration between observationally identified groups, which is in addition to the principal difficulty of deciding on the criteria and procedures that are used to group stars (as discussed at the beginning of \\kap{sec:final-ONa}). These complications impose a limit to the accuracy that we can expect in how well observed properties can be matched to the predictions of any scenario. The best that we can ask for at this point is rather qualitative agreement, and with this goal in mind we have summarized the alignment of the Galactic plane passage gas purging model with some recently reported observational properties of \\ocen\\ in \\tab{tab:gppgpm}.  While most of the entries in \\tab{tab:gppgpm} are based on our discussion in the previous sections there is a noteworthy peculiarity in the most metal-poor group (``first'' generation) which should represent the genuine first generation stars. In \\kap{sec:final-ONa} we made the case that the observational data may support the case of the superposition of separate sub-populations forming the overall anti-correlation in \\ocen. It seems that, in fact, in this lowest [Fe/H] bin there are signs of an intrinisic anti-correlation with at most an unclear signature of n-capture enhancement. Groups \\#6 and \\#4 of \\citet{Gratton:2011kr} taken together display a significant scatter in the [Na/O] ratio, while the lowest-metallicity bin of \\citet{DOrazi:2011jf} includes two very N-rich stars \\citep[see also][]{marino:12}. This possible substructure in the most metal-poor population in \\ocen\\ cannot be explained by the Galcatic plane passage gas purging model, and may be a relic of the proto-cluster object."
    },
    "1208/1208.4992_arXiv.txt": {
        "abstract": "The CoGeNT and CRESST WIMP direct detection experiments have recently observed excesses of nuclear recoil events, while the DAMA/LIBRA experiment has a long standing annual modulation signal. It has been suggested that these excesses may be due to light mass, $m_{\\chi}\\sim5-10 \\, {\\rm GeV}$, WIMPs. The Earth's motion with respect to the Galactic rest frame leads to a directional dependence in the WIMP scattering rate, providing a powerful signal of the Galactic origin of any recoil excess. We investigate whether direct detection experiments with directional sensitivity have the potential to observe this anisotropic scattering rate with the elastically scattering light WIMPs proposed to explain the observed excesses. We find that the number of recoils required to detect an anisotropic signal from light WIMPs at $5\\sigma$ significance varies from 7 to more than 190 over the set of target nuclei and energy thresholds expected for directional detectors. Smaller numbers arise from configurations where the detector is only sensitive to recoils from the highest speed, and hence most anisotropic, WIMPs. However, the event rate above threshold is very small in these cases, leading to the need for large experimental exposures to accumulate even a small number of events. To account for this sensitivity to the tail of the WIMP velocity distribution, whose shape is not well known, we consider two exemplar halo models spanning the range of possibilities. We also note that for an accurate calculation the Earth's orbital speed must be averaged over. We find that the exposures required to detect $10 \\, {\\rm GeV}$ WIMPs at a WIMP-proton cross-section of $10^{-4} \\, {\\rm pb}$ are of order $10^{3} \\, {\\rm kg \\, day}$ for a $20 \\, {\\rm keV}$ energy threshold, within reach of planned directional detectors. Lower WIMP masses require higher exposures and/or lower energy thresholds for detection. ",
        "introduction": "Direct detection experiments aim to detect dark matter in the form of Weakly Interacting Massive Particles (WIMPs) via the nuclear recoils which occur when WIMPs scatter off target nuclei~\\cite{DD}. The sensitivity of these experiments has increased rapidly over the last few years, and they are probing the regions of WIMP mass-cross-section parameter space populated by the lightest neutralino in Supersymmetric extensions of the standard model (see e.g. Ref.~\\cite{theory}).  Event rate excesses and annual modulations in various direct detection experiments have prompted recent interest in light WIMPs. The DAMA (now DAMA/LIBRA) collaboration have, for more than a decade, observed an annual modulation of the event rate in their ${\\rm NaI}$ crystals~\\cite{dama}. This annual modulation is consistent with light ($m_{\\chi} \\sim 5-10$ GeV) WIMPs scattering off ${\\rm Na}$~\\cite{Bottino:2003cz,gglightwimps}. The CoGeNT experiment, after allowing for backgrounds with an exponential plus constant energy spectrum, find an excess of low energy events which is consistent with WIMPs with mass $m_{\\chi} \\approx 7-11$ GeV~\\cite{cogent1}. With a larger data set they have observed a 2.8 $\\sigma$ annual modulation~\\cite{cogent2}, with period and phase broadly consistent with the expectation for WIMPs~\\cite{amtheory}. The CRESST experiment has observed an excess of events in their ${\\rm CaWO}_{4}$ crystals above expectations from backgrounds~\\cite{cresst}. The excess is compatible with either WIMPs of mass $m_{\\chi} \\sim 25$ GeV scattering off tungsten predominantly, or lighter, $m_{\\chi} \\sim 10$ GeV, WIMPs scattering off oxygen and calcium. It appears that it is not possible to explain all of these signals in terms of a single conventional elastic-scattering WIMP, especially when the exclusion limits from the CDMS~\\cite{cdms}, XENON10~\\cite{xenon10} and XENON100~\\cite{xenon100} experiments and the CRESST commissioning data~\\cite{cresstcomm} data are taken into account~\\cite{ksz,khb,ox} (see also Refs.~\\cite{hk,cogentamp}). None the less it is still possible that some subset of the putative signals are due to elastic scattering light WIMPs.  The deployment of a ${\\rm NaI}$ detector at the South Pole has been proposed to directly test the DAMA annual modulation signal~\\cite{damatest}. The direction dependence of the scattering rate~\\cite{dirndep} provides another potentially powerful way of testing whether the observed excesses and annual modulations are due to elastic scattering light WIMPs. The amplitude of the directional signal is far larger than that of the annual modulation and hence the anisotropy of the WIMP induced nuclear recoils could be confirmed with a relatively small number of events~\\cite{copi:krauss,pap1}. Furthermore the angular dependence of the recoils (in particular the peak recoil rate in the direction opposite to the motion of the solar system, or for low energy recoils  a ring around this direction~\\cite{ring}) is extremely unlikely to be mimicked by backgrounds, and would allow unambiguous detection of WIMPs~\\cite{billard,pap5}. In this paper we investigate whether current and near future directional detectors would be able to detect elastic scattering light WIMPs. ",
        "conclusions": "\\label{res} For each of the combinations of WIMP mass, target nuclei, energy threshold and confidence level discussed in Sec.~\\ref{model} we calculate the number of events required to reject isotropy, $N_{\\rm iso}$, for vector and axial data (using the Rayleigh-Watson and Bingham statistics respectively). The number of events required for a $5\\sigma$ detection with the Rayleigh-Watson statistic varies from 7 to 58. For fixed WIMP and target mass, $N_{{\\rm iso}}$ decreases with increasing energy threshold, $E_{\\rm th}$, (since the recoils caused by high speed WIMPs in the tail of the distribution are more anisotropic), until the minimum speed required to cause a recoil of energy $E_{\\rm th}$, $v_{\\rm min}(E_{\\rm th})$, exceeds the maximum WIMP speed in the lab frame, $v_{\\chi}^{\\rm max}$. At this point the event rate is zero and no events can be detected. The same trend occurs for decreasing WIMP mass (with threshold energy and target mass fixed). For fixed threshold energy, $N_{{\\rm iso}}$ decreases with increasing target mass for $m_{\\chi}= 5 \\, {\\rm GeV}$, however for $m_{\\chi}=7.5$ and $10 \\, {\\rm GeV}$, $N_{{\\rm iso}}$ increases as the target mass number is increased from $A=3$ to $12$ before decreasing as $A$ is increased further. This is due to the variation of $v_{\\rm min}$ with target and WIMP mass shown in Fig.~\\ref{vminfig}.  For the Bingham statistic, which can be used with axial data, the number of events required for a $5\\sigma$  detection varies from 9 to more than 190~\\footnote{In a small number of cases, where a large number of events are required, we have only been able to place a lower limit on $N_{\\rm iso}$ due to computational time limitations.}. Both the number of events and its increase, relative to the number required for the Rayleigh-Watson statistic, is smallest for the configurations which are only sensitive to the highly anisotropic recoils from high speed WIMPs in the tail of the speed distribution.  \\begin{figure}[t] \\includegraphics[width=8.5cm]{expostargetRW5sigMax.eps} \\caption{The exposure required for a $5 \\sigma$ rejection of isotropy as a function of energy threshold, $E_{\\rm th}$, for  ${}^3{\\rm He}$ (top left), C (top right), F (bottom left) and S (bottom right) for $m_{\\chi}= 5, 7.5$ and $10 \\, {\\rm GeV}$ (crosses, stars and triangles respectively) for the Maxwellian $f(v)$ with a sharp cut-off at $v_{\\rm  esc}^{\\rm max}= 608 \\, {\\rm km \\, s}^{-1}$. Where a symbol is not displayed, the minimum WIMP speed corresponding to the energy threshold, $v_{\\rm min}(E_{\\rm th})$, for this WIMP and target mass combination exceeds the maximum WIMP speed in the lab frame, $v_{\\chi}^{\\rm max}$ and hence the event rate is zero.} \\label{expostargetRW5sigMax} \\end{figure}   If an experiment is only sensitive to high speed WIMPs, fewer events are required to reject isotropy, how- ever, the reduced event rate means that the exposure required to accumulate these events will be larger. We therefore use eq.~(\\ref{exposeq}) to calculate the exposure, ${\\cal E}$, (in ${\\rm kg \\, day}$) required to accumulate the required number of events for each case. As illustrated in Fig.~\\ref{drde}, for light WIMPs the differential event rate decreases rapidly with increasing energy, and therefore the event rate above the energy threshold, $R(>E_{\\rm th})$, plays a crucial role in determining the exposure required.     The exposure required to reject isotropy at $5 \\sigma$ using the Rayleigh-Watson statistic assuming a Maxwellian $f(v)$ with a sharp cut-off at $v_{\\rm  esc}^{\\rm max}= 608 \\, {\\rm km \\, s}^{-1}$ is shown for each configuration in Fig.~\\ref{expostargetRW5sigMax}. While $N_{{\\rm iso}}$ varies by less than an order of magnitude, because of the large variation in $R(>E_{\\rm th})$, the exposure varies by more than five orders of magnitude. Due to the rapid decrease of $R(>E_{\\rm th})$, the exposure increases sharply with increasing $E_{\\rm  th}$ for each WIMP and target nuclei mass combination. The factor by which the exposure increases, increases with both decreasing WIMP mass and increasing target nuclei mass (i.e. as the minimum WIMP speed to which the experiment is sensitive is increased). Eventually the minimum WIMP speed corresponding to the energy threshold, $v_{\\rm min}(E_{\\rm th})$,  exceeds the maximum WIMP speed in the lab frame, $v_{\\chi}^{\\rm max}$, and the event rate is zero and WIMPs of this mass can not be detected.   \\begin{figure} \\includegraphics[width=8.5cm]{expostargetRW5sigS.eps} \\caption{As Fig.~\\ref{expostargetRW5sigMax} but for a S target only comparing  the exposures required for the  Lisanti et al. $f(v)$, eq.~(\\ref{k}), with $v_{\\rm esc}^{\\rm min}= 498 \\, {\\rm km \\, s}^{-1}$ (upper symbols) with those for the Maxwellian $f(v)$ with a sharp cut-off at $v_{\\rm  esc}^{\\rm max}= 608 \\, {\\rm km \\, s}^{-1}$ (lower symbols). For $m_{\\chi} = 7.5 \\, {\\rm keV}$ and $E_{\\rm th}= 5 \\, {\\rm keV}$ the exposures required for the two speed distributions are the same, and hence only one symbol is visible.} \\label{expostargetRW5sigS} \\end{figure}    Of the halo models considered, the Maxwellian $f(v)$ with a sharp cut-off at $v_{\\rm  esc}^{\\rm max}= 608 \\, {\\rm km \\, s}^{-1}$ has the largest tail event rate, and hence the smallest exposures. In Fig.~\\ref{expostargetRW5sigS} we show the exposures for a ${\\rm S}$ target for the Lisanti et al. $f(v)$, eq.~(\\ref{k}), with $v_{\\rm esc}^{\\rm min}= 498 \\,{\\rm km \\, s}^{-1}$ as well. When $v_{\\rm min}(E_{\\rm th})$ is much smaller than $v_{\\chi}^{\\rm max}$ the exposure required is fairly modest, ${\\cal E} \\sim 10-100 \\, {\\rm kg} \\, {\\rm day}$ and the event rates, and hence exposures, for the two speed distributions are very similar. However as $v_{\\rm min}(E_{\\rm th})$ approaches $v_{\\chi}^{\\rm max}$ the exposures required become large and the differences between the two speed distributions become significant. For instance for $E_{\\rm th} = 20 \\, {\\rm keV}$, WIMPs with $m_{\\chi} =10 \\, {\\rm GeV}$ and a Maxwellian distribution with $v_{\\rm  esc}^{\\rm max}= 608 \\, {\\rm km \\, s}^{-1}$ anisotropy could be detected with an exposure of $1700  \\, {\\rm kg} \\, {\\rm day}$, however if the WIMPs have the Lisanti et al. $f(v)$ with $v_{\\rm esc}^{\\rm min}= 498 \\, {\\rm km \\, s}^{-1}$ the event rate is zero and they can not be detected. The trends for the other target nuclei are similar.  In Fig.~\\ref{exposRWB5sigS} we compare the exposures required to reject isotropy with axial data using the Bingham statistic with those for vectorial data using the Rayleigh-Watson statistic, for a ${\\rm S}$ target and a Maxwellian $f(v)$ with a sharp cut-off at $v_{\\rm  esc}^{\\rm max}= 608 \\, {\\rm km \\, s}^{-1}$. Since $R(>E_{\\rm th})$ for each configuration doesn't change, the variations in the exposure are driven entirely by the variations in $N_{\\rm iso}$ discussed above. Therefore the increase in the exposure, relative to that required for the Rayleigh-Watson statistic, is smallest for the cases which are only sensitive to the highly anisotropic recoils from high speed WIMPs in the tail of the speed distribution. However in these cases the exposure is large even for the Rayleigh-Watson statistic, due to the small values of $R(>E_{\\rm th})$.  \\begin{figure} \\includegraphics[width=8.5cm]{exposRWB5sigS.eps} \\caption{As Fig.~\\ref{expostargetRW5sigS} but comparing  the exposures required using the Bingham statistic (upper symbols) with those required using the Rayleigh statistic (lower symbols), in both cases for a S target using a Maxwellian $f(v)$ with a sharp cut-off at $v_{\\rm  esc}^{\\rm max}= 608 \\, {\\rm km \\, s}^{-1}$. For $m_{\\chi} = 10 \\, {\\rm GeV}$ and $E_{\\rm th}=5 \\, {\\rm keV}$ we have only been able to place a lower bound on the number of events, and hence exposure, required with the Bingham statistic. } \\label{exposRWB5sigS} \\end{figure}    We have focused on the number of events and exposure required for a `$5 \\sigma$' discovery of light WIMPs with a directional detector. Significant experimental support for light WIMPs would be obtained even with a lower significance signal. For $95\\%$ and $99.713\\%$ confidence levels (the later corresponding to $3 \\sigma$) the number of events, and hence the exposure, required with the Rayleigh-Watson statistic is smaller by a factor of between $0.40-0.91$ and $0.69-0.95$ respectively. The factor by which the number of events must be increased to increase the significance of the rejection of isotropy is smallest when the experiment is only sensitive to the most anisotropic events coming from the tail of the speed distribution. However in these cases the exposure required even for a low significance detection is large.    The final question is `How achievable are these exposures and energy thresholds by current and near-future detectors'?  A typical current detector consisting of a $ 1 \\, {\\rm m}^{3} $ TPC filled to 75 Torr with ${\\rm CF}_{4}$ or ${\\rm  CS}_{2}$ could, in roughly a year, achieve an exposure of order $10^{3} \\, {\\rm kg \\, day}$~\\cite{cygnus}. A $10^{3} \\, {\\rm kg \\, day}$ exposure with ${\\rm CF}_{4}$ or ${\\rm  CS}_{2}$ would (provided that recoils are measured in 3d with good, $\\lesssim 10^{\\circ}$, angular resolution) be capable of detecting $m_{\\chi}=10 \\, {\\rm keV}$ WIMPs with an energy threshold of $20 \\, {\\rm keV}$ or lower. Lighter WIMPs would require a lower energy threshold, potentially as low as $5 \\,{\\rm keV}$ for $m_{\\chi}= 5 \\, {\\rm GeV}$. With a ${}^3{\\rm He}$ target a low, $\\sim 5 \\, {\\rm keV}$, energy threshold would be required, even for $m_{\\chi}= 10 \\, {\\rm GeV}$. Measuring the directions of low energy nuclear recoils is a major experimental challenge, and these energy threshold are lower than those which have been used in the analysis of data from current, prototype, detectors~\\cite{cygnus}. One of the focuses of the R\\&D for future generation experiments is to reduce the energy threshold (see e.g. \\cite{cygtalks}). The MIMAC experiment has, using micromegas readout, detected 5 keV F recoils~\\cite{mimaclow}, while DRIFT-II is sensitive to nuclear recoils down to sub $5$ keV energies~\\cite{driftlow}.  Simulations of the MIMAC detector indicate that the directional energy threshold will lie below $20$ keV~\\cite{billardtrack}"
    },
    "1208/1208.4489_arXiv.txt": {
        "abstract": "We investigated the ALMA science verification data of Orion~KL and found a spectral signature of the vibrationally excited H$_{2}$O maser line at 232.68670 GHz ($\\nu_{2}$=1, 5$_{5,0}$$-$6$_{4,3}$). This line has been detected in circumstellar envelopes of late-type stars so far but not in young stellar objects including Orion~KL. Thus, this is the first detection of the 232~GHz vibrationally excited H$_{2}$O maser in star-forming regions. The distribution of the 232~GHz maser is concentrated at the position of the radio Source~I, which is remarkably different from other molecular lines. The spectrum shows a double-peak structure at the peak velocities of $-$2.1 and 13.3~km~s$^{-1}$. It appears to be consistent with the 22~GHz H$_{2}$O masers and 43~GHz SiO masers observed around Source~I. Thus, the 232~GHz H$_{2}$O maser around Source~I would be excited by the internal heating by an embedded protostar, being associated with either the root of the outflows/jets or the circumstellar disk around Source~I, as traced by the 22~GHz H$_{2}$O masers or 43~GHz SiO masers, respectively. ",
        "introduction": "Water is one of the most abundant interstellar molecules after H$_{2}$ and hence, it is important for interstellar chemistry and physics of molecular clouds \\citep[e.g.][]{vandishoeck2011}. However, due to the large atmospheric opacity, ground-based observations of the H$_{2}$O lines in radio and infrared bands are almost impossible except for the isotopic species (e.g. HDO and H$_{2}^{18}$O) and strong maser lines. In particular, the 6$_{1, 6}$$-$5$_{2, 3}$ transition at 22~GHz (lower state energy, $E_{l}$=642~K) is known to show extremely strong maser emission in circumstellar envelopes (CSEs) around late-type stars, young stellar objects (YSOs) in star-forming regions (SFRs), and active galactic nuclei. The 22~GHz maser has been used as a unique probe of dense gas and their dynamics with very long baseline interferometers (VLBI) thanks to its extremely high brightness and compact structure \\citep[e.g.][]{chapman2007}.  Other H$_{2}$O maser lines are also detected in millimeter/submillimeter wavelength \\citep{humphreys2007}. Lower excitation lines in 183~GHz ($E_{l}$=196~K) and 325~GHz ($E_{l}$=454~K) are detected both in CSEs and around YSOs while some of the higher excitation lines including vibrationally excited lines are detected only in CSEs. Multi-transition studies of H$_{2}$O maser lines could be powerful tools to investigate shocked regions in CSEs and YSOs at the highest spatial resolution achieved with VLBI and millimeter/submillimeter interferometers when combined with the theoretical models \\citep{neufeld1990, neufeld1991}.  In this Letter, we report the detection of the vibrationally excited H$_{2}$O maser line at 232.68670~GHz ($E_{l}$=3451~K) in a massive SFR Orion~KL at a distance of 420~pc \\citep{hirota2007, kim2008} with the Atacama Large Millimeter/Submillimeter Array (ALMA). This line has been detected in late-type stars so far but not in YSOs including Orion~KL \\citep{menten1989}. Thus, this is the first detection of the 232~GHz H$_{2}$O maser line in YSOs. ",
        "conclusions": "As discussed above, we can identify the 232.68670~GHz feature detected in the ALMA SV data for Orion~KL as the vibrationally excited H$_{2}$O maser. This is the first detection of this maser line in YSOs. The peak flux of the {\\it{blended}} feature observed with the 2\\arcsec$\\times$2\\arcsec \\ aperture is 0.43~Jy (Figures \\ref{fig-spectrum} and \\ref{fig-maser}). It corresponds to the brightness temperature of 2.4~K. If this emitting region is as compact as this aperture size, the single-dish observation with the beam size of 30\\arcsec \\ would yield the brightness temperature of 0.01~K. Thus, it was not detectable with the previous observations with single-dish telescopes \\citep{menten1989}, although a possible spectral feature can be seen in the line survey data by \\citet{sutton1985}, probably attributed to the HCOOCH$_{3}$ line.  The vibrationally excited H$_{2}$O masers have been detected in only several oxygen-rich late-type stars \\citep{menten1989}, which could be attributed to their higher excitation levels (3451~K) than that of the 22~GHz maser (642~K). The 232~GHz H$_{2}$O maser around Source~I would be excited due to the internal heating by an embedded YSO as expected from a maser pumping mechanism for late-type stars. Source~I is also known as a powering source of the SiO masers, which is quite rare for YSOs \\citep{zapata2009}. This observational evidence may imply similar characteristics of Source~I and late-type stars. Further studies with the millimeter/submillimeter masers in Source~I, along with other YSOs and CSEs in late-type stars, will be crucial in understanding the pumping mechanism of the H$_{2}$O maser lines, physical and dynamical state of these maser sources, and accordingly mass-loss/accretion processes occurring in the YSOs and CSEs.  The 232.68670~GHz H$_{2}$O maser emission is concentrated around Source~I. However, the distribution of the 232~GHz H$_{2}$O maser features could not be resolved with the ALMA SV data with the beam size of 1\\arcsec.7$\\times$1\\arcsec.4. According to the double-peaked spectra of the 22~GHz and 232~GHz H$_{2}$O maser as shown in Figure \\ref{fig-maser}, the 232~GHz maser features would have similar structure to that of the 22~GHz masers rather than the SiO masers. Higher resolution imaging will reveal their spatial structure and provide information about a  possible powering source of the 232~GHz H$_{2}$O maser; whether they are really associated with the root of outflows/jets as traced by the 22~GHz H$_{2}$O masers \\citep{gaume1998} or with circumstellar disk as traced by the 43~GHz SiO masers \\citep{reid2007, kim2008}.  In the present study, we could not perfectly separate the contribution from the HCOOCH$_{3}$ and the 232~GHz H$_{2}$O maser lines mainly due to the insufficient spatial resolution. Therefore, it is still unclear whether the 232~GHz masers are distributed other than in Source~I, such as the Compact Ridge where strong H$_{2}$O maser lines are sometimes detected \\citep{hirota2011, gaume1998}. A search for the 232.68670~GHz H$_{2}$O maser lines with higher spatial resolution would be important to distinguish their distribution by filtering out the contribution from the thermal and extended HCOOCH$_{3}$ emission.  Our results clearly demonstrate ALMA's high possibility of detecting new millimeter/submillimeter maser lines with its high sensitivity/resolution. Further observational studies with ALMA of millimeter/submillimeter masers will reveal spatial and velocity structure of the maser sources at a resolution of $\\sim$10~mas or better, providing a complementary method to the VLBI studies of the SiO and 22~GHz H$_{2}$O masers."
    },
    "1208/1208.3693_arXiv.txt": {
        "abstract": "{An object with a very peculiar light-curve was discovered recently using Kepler data. Authors argue that this object may be a transiting disintegrating planet with a comet like dusty tail. Light-curves of some eclipsing binaries may have features analogous to this light-curve and it is very interesting to see whether they are also caused by the same effects and put them in a more general context. } {The aim of the present paper is to verify the model suggested by the discoverers by the light-curve modelling and put constraints on the geometry of the dust region and various dust properties.} {We modify the code SHELLSPEC designed for modelling of the interacting binaries to calculate the light-curves of stars with such planets. We take into account the Mie absorption and scattering on spherical dust grains of various sizes assuming realistic dust opacities and phase functions and finite radius of the source of the scattered light. } { The planet light-curve is reanalysed using long and short cadence Kepler observations from the first 14 quarters. Orbital period of the planet was improved. We prove that the peculiar light-curve of this objects is in agreement with the idea of a planet with a comet like tail. Light-curve has a prominent pre-transit brightening and a less prominent post-transit brightening. Both are caused by the forward scattering and are a strong function of the particle size. This feature enabled us to estimate a typical particle size (radius) in the dust tail of about 0.1-1 micron. However, there is an indication that the particle size changes along the tail. Larger particles better reproduce the pre-transit brightening and transit core while smaller particles are more compatible with the egress and post-transit brightening. Dust density in the tail is a steep decreasing function of the distance from the planet which indicates a significant tail destruction caused by the star. We also argue that the 'planet' does not show uniform behaviour but may have at least two constituents. This light-curve with pre-transit brightening is analogous to the light-curve of $\\epsilon$ Aur with mid-eclipse brightening and forward scattering plays a significant role in such eclipsing systems. } {} ",
        "introduction": "The exoplanet candidate KIC012557548b has been discovered recently by Rappaport et al. (\\cite{rappaport12}). It was discovered from Kepler long cadence data (Borucki et al. \\cite{borucki11}) obtained during first two quarters. This exoplanet is very unique. Unlike all other exoplanets it exhibits strong variability in the transit depth. For some period of time transits even disappear. Shape of the transit is highly asymmetric with a significant brightening just before the eclipse, sharp ingress followed by a smooth egress. Planet has also extremely short period of 0.65356(1) days (15.6854 hours). Rappaport et al. (\\cite{rappaport12}) suggested that the planet has size not larger than Mercury and is slowly disintegrating/evaporating what creates a comet like tail made of pyroxene grains. Perez-Becker \\& Chiang \\cite{perez13} constructed a radiative-hydrodynamic model of the atmospheric escape from such low mass rocky planets. The hypothesis that a close-in planet can have a cometary-like tail was first suggested by Schneider et al. (\\cite{schneider98}) and revisited by Mura et al. (\\cite{mura11}). The transit light-curve of dusty comets was first investigated by Lecavelier des Etangs et al. (\\cite{lecavelier99}). There is another class of objects which may look very different but may have features analogous to this light-curve. $\\epsilon$ Aur is an interacting binary with the longest known orbital period, 27.1 yr. The primary star, which is the main source of light, may be either a young massive F0Ia super-giant or an evolved post-AGB star (see Guinan et al. \\cite{guinan12}, Hoard et al. \\cite{hoard} and references therein). The star is partially eclipsed by a dark dusty disk and the light-curve shows a very unusual shallow mid-eclipse brightening (MEB). It was suggested that the disk has a central hole and is inclined out of the orbital plane so that the star can peek through the hole (Carroll et al. \\cite{carroll}). Recently, Budaj (\\cite{budaj11a}) proposed that MEB might be due to the flared disk geometry and forward scattering on dust and that one does not necessarily have to see the primary star through the hole in the disk during the eclipse. Calculation of Muthumariappan \\& Parthasarathy (\\cite{muthumariappan12}) indeed indicate that the disk is flared and the hole is present but not seen at such almost edge-on inclination. Shallow MEB was observed also in some other long period eclipsing binaries, for example in AZ Cas (Galan et al. \\cite{galan12}).  Planetary transits are usually modelled using the analytical formulae of Mandel \\& Agol (\\cite{mandel02}). This approach assumes spherical shape of the objects. JKTEBOP code can calculate and solve the transits numerically assuming the shape of bi-axial ellipsoids (Southworth \\cite{southworth12}). BEER algorithm (Faigler \\& Mazeh \\cite{faigler11}) calculates analytical light-curves including Doppler boosting. The EVILMC code developed by Jackson et al. (\\cite{jackson12}) can model transit light-curves assuming the Roche shape. The SHELLSPEC code of Budaj \\& Richards (\\cite{budaj04}) can calculate planetary light-curves assuming the Roche shape, the new model of the reflection effect, and the circum-stellar/planetary material. Modelling the transit light-curve of KIC012557548, however, will be different and will require at least some modifications to these codes or a new approach. Shortly before the submission of this manuscript the light-curve of this planet was modelled independently by Brogi et al. (\\cite{brogi12}). These authors started with thorough data reduction of raw Kepler long cadence photometry from first six quarters. They assumed 1D model of the dust cloud in which the vertical dimension of the cloud was negligible compared with the stellar radius. They also modelled dust extinction as a free parameter and assumed analytical Henyey-Geenstein phase functions and point source approximation for the scattered light. The particle size was estimated mainly from the overall shape of the transit.  In this paper we first revisit the light-curve and orbital period using long as well as short cadence Kepler observations from the first 14 quadratures (Sections \\ref{obs}, \\ref{tail}). Then we modify the code SHELLSPEC to model light-curves of such objects. In Section \\ref{dust} we will calculate real opacities and phase functions of pyroxene and other similar dust grains. In Section \\ref{s4} we construct a 3D model of the dust cloud, calculate the radiative transfer along the line of sight and take into account a finite dimension of the source of light. We will estimate the particle size from the pre-transit brightening feature which is most sensitive to this parameter. This will enable us to verify whether the shape of the light-curve is in agreement with the idea of a planet with the comet like tail, put constraints on the particle size and geometry of the dust region, and put this interesting object which may look like a comet, behave like an eclipsing interacting binary, but be an exoplanet, into a more general context.  \\begin{figure} \\centering \\includegraphics[angle=0,width=8.3cm]{beta_scan.eps} \\caption{ The search for a long term period change. The minimum of the theta parameter corresponds to the period change rate $\\beta=0.3\\pm0.5$ days/Myr which means that there is no significant long term period variability. } \\label{f1} \\end{figure} ",
        "conclusions": "The light-curve of this planet candidate was reanalysed using first 14 quarters of the Kepler data including first short cadence observations.  Orbital period of the planet was improved. We searched for the long term period changes but found no convincing evidence of such changes.  Quasi periodic variability in the tail of the planet with the period of about 1.5 year was discovered.  We modelled the light-curve of KIC012557548 using SHELLSPEC code with the following assumptions: spherical dust grains, different dust species (pyroxene, enstatite, forsterite), different particle sizes, Mie absorption and scattering, finite radius of the source of light (star) and that the medium is optically thin. We proved that its peculiar light-curve is in agreement with the idea of a planet with a comet like tail. A model with a dusty ring in the orbital plane which has an inclination of about 82 degrees, exponential or power law density profile ($A1=A2=-20$) fits the observations surprisingly well.  We confirmed that the light-curve has a prominent pre-transit brightening. There is an indication of a less prominent brightening after the transit, both are caused by the forward scattering.  Dust density in the tail is a steep decreasing function of the distance from the planet which indicates significant destruction of the tail caused by the star.  Transit depth is a highly degenerate function of the particle size, dust density, and other dust properties. Various combinations of them were estimated and tabulated. There will also be a degeneracy between the inclination, particle size, and the density exponent.  However, the forward scattering and the pre(post)-transit brightening are quite sensitive to the particle size. Consequently, we estimated the particle size (radius) of the grains in the head of the tail from the pre-transit brightening to be about 0.1-1 micron. There is an indication that the particle size is larger at the head and decreases along the tail to about 0.01-0.1 micron.  We argue that there are several indications that the 'planet' is not homogeneous and that it consists of several components. The component that is responsible for the transit core and the other responsible for the tail. Components may have different grains with different density profiles, and/or particle size.  It is interesting to note that this planet's light-curve with pre-transit brightening is analogous to light-curves of some interacting binaries with mid-eclipse brightening, particularly $\\epsilon$ Aur and forward scattering plays an important role in all of them."
    },
    "1208/1208.1696_arXiv.txt": {
        "abstract": "{Detecting cosmic ray hits (cosmics) in fiber-fed integral-field spectroscopy (IFS) data of single exposures is a challenging task because of the complex signal recorded by IFS instruments. Existing detection algorithms are commonly found to be unreliable in the case of IFS data, and the optimal parameter settings are usually unknown apriori for a given dataset. } {The Calar Alto legacy integral field area (CALIFA) survey generates hundreds of IFS datasets for which a reliable and robust detection algorithm for cosmics is required as an important part of the fully automatic CALIFA data reduction pipeline. Such a new algorithm needs to be  tested against the  performance of the commonly used algorithms \\texttt{L.A.Cosmic} and \\texttt{DCR}. General recommendations for the usage and optimal parameter settings of each algorithm have not yet been systematically studied for fiber-fed IFS datasets to guide users in their choice.} {We developed a novel algorithm, \\texttt{PyCosmic}, which combines the edge-detection algorithm of \\texttt{L.A.Cosmic} with a point-spread function convolution scheme. We generated mock data to compute the efficiency of different algorithms for a wide range of characteristic fiber-fed IFS datasets using the Potsdam Multi-Aperture Spectrophotometer (PMAS) and the VIsible MultiObject Spectrograph (VIMOS) IFS instruments as representative cases.} {\\texttt{PyCosmic} is the only algorithm that achieves an acceptable detection performance for CALIFA data. We find that \\texttt{PyCosmic} is the most robust tool with a detection rate of $\\gtrsim90$\\% and a false detection rate $\\lesssim5$\\% for any of the tested IFS data. It has one less free parameter than the \\texttt{L.A.Cosmic} algorithm. Only for strongly undersampled IFS data does \\texttt{L.A.Cosmic} exceed the performance of \\texttt{PyCosmic} by a few per cent. \\texttt{DCR} never reaches the efficiency of the other two algorithms and should only be used if computational speed is a concern. Thus, \\texttt{PyCosmic} appears to be the most versatile cosmics detection algorithm for IFS data. It is implemented in the new CALIFA data reduction pipeline as well as in recent versions of the multi-instrument IFS pipeline \\texttt{P3D}. Although \\texttt{PyCosmic} has been optimized for IFS data, we have also successfully applied it to longslit data and anticipate that good results will be achieved with imaging data.} {} ",
        "introduction": "The identification and rejection of artefacts on charged-couple device (CCD) detectors caused by cosmic ray hits (hereafter cosmics) is a persisting problem for the reduction and analysis of astronomical data. Combining multiple images of the same object or field is considered the best method to identify cosmics because it is less likely that the same pixel is affected in several images. Sophisticated algorithms that detect and reject outlier pixels during the combination of exposures were developed, for example, for the Hubble Space Telescope \\citep[e.g.,][]{Fruchter:1997}. However, there are often cases where only a single exposure is available or multiple exposures cannot be combined. This happens frequently with fiber-fed integral field spectroscopic (IFS) data, where the effects of differential atmospheric refraction, instrument flexure, or a variable sky brightness during the sequence of exposures prevent reliable detection of cosmics by image comparison.  Various techniques have been developed to detect and reject cosmics in single CCD exposures. They use different methods like trained neural networks \\citep{Salzberg:1995}, convolution with a point-spread function \\citep[PSF,][]{Rhoads:2000}, Laplacian edge detection \\citep[][hereafter D01]{Dokkum:2001}, image statistics \\citep[][hereafter P04]{Pych:2004}, or a fuzzy logic approach \\citep{Shamir:2005}. A detailed performance evaluation of the different algorithms on single astronomical images was presented by \\citet{Farage:2005}. Their tests revealed that the D01 algorithm, also known as \\texttt{L.A.Cosmic}, performed well on imaging data. The algorithm of P04, known as \\texttt{DCR}, did not perform as well on images, but was much less computationally expensive and primarily designed for spectroscopic data.  Currently, a thorough evaluation of the performance of cosmics detection algorithms for fiber-fed IFS data is missing. Signals in such data are much more complex because a spectrum from each individual fiber is recorded  along a discrete trace on the CCD, with little gaps between spectra.  Thus, edge-like structures are introduced, and bright object or night-sky emission lines are more likely to be misclassified as cosmics, which is why automatic data-reduction pipelines generally avoid including this crucial step in the reduction process  \\citep[e.g.,][]{Barnsley:2012}. Sophisticated methods to detect cosmics in data from fiber-fed multi-object spectrographs were presented \\citep{Zhu:2009,Wang:2009} and show excellent results, but their parameter choices seem arbitrary for the \\texttt{L.A.Cosmic} and \\texttt{DCR} algorithms as their prime reference. Additionally, there is no public code available to make an independent check of their results and to verify whether the algorithm works also with IFS data.  For the Calar Alto legacy integral field area (CALIFA) survey \\citep{Sanchez:2012a} and other IFS studies using the same instrument \\citep[e.g.,][]{Sandin:2008}, it was discovered that the available algorithms always selected night-sky or object-emission lines as cosmics for the IFS data. An initial attempt to reduce the high false detection rate for CALIFA data by using a simplified Laplacian edge detection algorithm was implemented into the \\texttt{R3D} reduction package \\citep{Sanchez:2006a} and was only partly successful. Although it reduced the number of false detections, a significant number of cosmics were undetected.  In this paper, we present a novel algorithm called \\texttt{PyCosmic}. It combines the iterative Laplacian edge detection scheme with a PSF convolution approach. We evaluate the performance of \\texttt{PyCosmic} against the most popular algorithms available, \\texttt{DCR} and \\texttt{L.A.Cosmic}, on realistic mock data for different IFSs and compare the results with illustrative examples on observed raw data. We then provide general recommendations regarding the use of detection algorithms with fiber-fed IFS data.  The different algorithms used in this study are briefly described in Section~\\ref{sec:methods}. Results of our detailed performance and parameter study on IFS mock data are then presented in Section~\\ref{sect:parameter}, followed by results obtained for real data in Section~\\ref{sect:real}. Finally, we provide general recommendations as a guide for other IFS users in Section~\\ref{sect:guide}. ",
        "conclusions": "In this paper we have presented a novel detection algorithm for cosmics in single exposures called \\texttt{PyCosmic}. The algorithm combines Laplacian edge detection with a PSF convolution approach. We systematically compared the performance of our new algorithm  against other standard detection algorithms, \\texttt{DCR} \\citep{Pych:2004} and \\texttt{L.A.Cosmic} \\citep{Dokkum:2001}, on simulated and real images from fiber-fed IFS instruments.  With the aid of these detailed comparison tests, we provide general recommendations for the use of these algorithms for the detection of cosmics in IFS data.  We have found that \\texttt{DCR} does not reach a detection efficiency equivalent to that of \\texttt{L.A.Cosmic} and \\texttt{PyCosmic}. Therefore, we cannot recommend its use for IFS data in general, except when computational speed is critical. The strength of the \\texttt{L.A.Cosmic} algorithm is that it works best for undersampled IFS data. However, a significant drawback is that the minimum false detection rate achievable for a given IFS data is entirely set by the characteristics of the instrument and cannot be reduced by changing any parameter settings. This peculiarity is most evident for PMAS-PPak IFU data from the CALIFA survey \\citep{Sanchez:2012a}, where the false detection rate of \\texttt{L.A.Cosmic} is $P_\\mathrm{f}\\gtrsim40\\%$.  Our \\texttt{PyCosmic} algorithm reduces the false detection rate with different parameter settings and solves this problem effectively. It has replaced the simplified \\texttt{R3D} routine (based solely on a Laplacian edge detection scheme) in the reduction pipeline of the CALIFA survey.  \\texttt{PyCosmic} is the most robust detection algorithm for cosmics in fiber-fed IFS data. In combination with well-characterized optimal parameter settings, it is well-suited for automatic usage for very large datasets. CALIFA is already a huge IFS survey by current standards that has significantly benefited from the development of \\texttt{PyCosmic}. The next generation of IFS instruments like the Sydney-AAO Multi-object IFS \\citep[SAMI,][]{Croom:2012} or the IFU project Mapping Nearby Galaxies at APO (MaNGA) is already being built or is planned to carry out even larger IFS surveys. These surveys will deliver IFS data for thousands of galaxies in the near future, which will certainly benefit from robust data reduction algorithms such as \\texttt{PyCosmic}.  The \\texttt{PyCosmic} algorithm has recently been implemented in the versatile multi-IFU reduction software \\texttt{P3D} \\citep{Sandin:2010} and is also available as a \\texttt{Python}-based stand-alone program\\footnote{\\texttt{PyCosmic} is available for download at \\url{http://pycosmic.sf.net}} so that it can be easily used or even added to any existing IFS reduction pipeline. Although \\texttt{PyCosmic} has been optimized for IFS data, we have also applied it successfully to longslit data and anticipate that good results will be achieved with imaging data."
    },
    "1208/1208.4599_arXiv.txt": {
        "abstract": "We present results for the QUEST RR Lyrae Survey at low galactic latitude, conducted entirely with observations obtained with the QUEST mosaic camera and the 1.0/1.5m J\\\"urgen Stock Schmidt telescope at the National Observatory of Venezuela. The survey spans an area of $476$ deg$^2$ on the sky, with multi-epoch observations in the $V$, $R$, and $I$ photometric bands for $6.5\\times10^6$ stars in the galactic latitude range $-30\\degr \\leqslant b \\leqslant +25\\degr$, in a direction close to the Galactic Anticenter $190\\degr \\leqslant l \\leqslant 230\\degr$. The variability survey has a typical number of $30$ observations per object in $V$ and $I$  and $\\sim25$ in $R$, with up to $\\sim120-150$ epochs in $V$ and $I$ and up to $\\sim100$ in $R$ in the best sampled regions. The completeness magnitudes of the survey are $V=R=18.5$ mag, and $I=18.0$ mag. We  identified 211 RR Lyrae stars, 160 \\textit{bona fide} stars of type \\typeab~and 51 candidates of type~\\typec, ours being the first \\emph{deep} RR Lyrae survey conducted at low galactic latitude.The completeness of the RR Lyrae survey was estimated in $\\ga95$ per cent and $\\sim85$ per cent for \\rrab~and  \\rrc~stars respectively. Photometric metallicities were computed based on the light curves and individual extinctions calculated from minimum light colours for each \\rrab~star. Distances were obtained with typical errors $\\sim7$ per cent. The RR Lyrae survey simultaneously spans a large range of heliocentric distances $0.5 \\leqslant R_{hel}\\mathrm{(kpc)} \\leqslant 40$ and heights above the plane $-15\\leqslant z\\mathrm{(kpc)} \\leqslant +20$, with well known completeness across the survey area, making it an ideal set for studying the structure of the Galactic thick disk. ",
        "introduction": "\\label{s:intro}  In the Milky Way (MW), the thick disk hosts a very old ($>10$ Gyr) and relatively metal-poor ($\\FeH\\sim-0.7$) stellar population \\citep{Wyse2009,Reddy2008}, comprising $\\sim10$ per cent or even up to $\\sim20$ of the thin disk mass \\citep[e.g.][]{Juric2008}. The Galactic thick disk therefore constitutes an important fossil record of the earliest stages of the formation of the Galactic disk, which could help understand the relevant mechanisms contributing to the formation of our Galaxy. From the perspective of galaxy formation, thick disks are remarkably relevant since these have proven to be an ubiquitous component of disk galaxies, as external galaxy surveys have shown that approximately $95$ per cent of disk galaxies contain thick disks \\citep{Yoachim2006}.  \\begin{figure*} \\begin{center} \\includegraphics[width=1.9\\columnwidth]{figure1.eps} \\caption{Aitoff projection map showing the footprints of recent large-scale RRLS surveys in Galactic coordinates. \\emph{Left:} Deep surveys ($V>16$), \\emph{(blue)} Present survey, \\emph{(cyan)} QUEST halo RRLS survey, \\emph{(orange)} SDSS Stripe 82, \\emph{(green)} LONEOS,\\emph{(red)} SEKBO, \\emph{(purple)} MACHO, \\emph{(dark yellow)} OGLE. \\emph{Right:} Shallow surveys ($V<16$), \\emph{(magenta)} NSVS, \\emph{(yellow)} ASAS-3 and \\emph{(light green)} HATNET. The solid black line corresponds to the Celestial Equator. Survey characteristics are summarized in Table \\ref{t:survey_details}.} \\label{f:survey_coverage_aitoff} \\end{center} \\end{figure*}   On the other hand, despite many advances in the last few years, there is no consensus about the formation mechanism of the Galactic thick disk, or even on some fundamental properties of its structure, mainly due to observational difficulties. Scale length measurements for the MW thick disk range from $h_R\\sim2$ to $4.7$ kpc \\citep{Carollo2010,Larsen2003,CabreraLavers2005,Chiba2000}, which makes it unclear whether the thick disk is more radially extended than the thin disk or not, or how the density profile behaves at large radii. Another example are metal abundances, the iron abundance of thick disk stars seems to be fairly homogeneous in the vertical direction \\citep[e.g.][]{Katz2011,Soubiran2003}, although \\citet{Chen2011} recently report an appreciable metallicity gradient. Radial metallicity gradients have not been probed yet and there is also ongoing debate regarding the reach in $\\FeH$ at both the low and high metallicity tails of the thick disk \\citep{Bensby2007,Reddy2008,Reddy2005}. In a recent paper, \\citet{Bovy2011,Bovy2011b} argue against the need of a thin/thick disk decomposition, proposing that the Galactic disk can be described as a series of simple stellar populations having scale lengths and scale heights which vary smoothly as a function of metallicity and $\\alpha-$element abundance. This illustrates how the issue of Galactic thick disk structure is far from settled and reliable measurements on structural, kinematic and chemical properties are necessary to gain a better understanding of the MW's thick disk which in turn, will impose strong constraints on the possible contribution of the different formation mechanisms.  Observationally, disentangling the structure of the Galactic disk is a challenging task since the Sun is embedded in it. Numerous studies have been conducted at high galactic latitudes ($|b|\\ga30\\degr$) in order to study stars at a height above the Galactic plane $z\\sim1-3$ kpc such that the thin disk density has decayed sufficiently, reducing contamination \\citep[e.g.][]{Ojha2001,Siegel2002,Brown2008b,Wyse2009}, but also limiting the coverage in the radial direction. At intermediate to low latitudes ($|b|\\la30\\degr$) the problems of crowding and the high and variable extinction become important, and the contribution of the thin disk has to be modelled \\citep{Robin1996,Larsen2003,Carollo2010,deJong2010}.  RR Lyrae stars (RRLS) have several fundamental properties that offer advantages in dealing with these problems inherent to low galactic latitude observations. RRLSs are \\emph{luminosity and colour} standards. Since RRLSs are horizontal branch stars, they have a well known absolute magnitude which shows little spread \\citep{Smith1995}, and allows the computation of distances with small uncertainties \\citep[$\\la10$ per cent,][hereafter \\citetalias{Vivas2004}]{Vivas2004}. Aditionally, during the phase of minimum light of the pulsation cycle, RRLSs of type \\typeab~have approximately the same effective temperature, and thus show very little dispersion in colour \\citep{Sturch1966,Day2002}. This property enables the use of \\rrab~stars as colour standards to measure extinctions up to the distance of each individual star; which is a crucial point in a low latitude survey since reddenning changes drastically along different lines of sight and has a strong dependence with distance. Also, since RRLSs trace old ($\\ga$10 Gyr) and mainly metal-poor ($\\FeH<-0.5$) stellar populations \\citep{Smith1995,Demarque2000}, the expected contamination from the thin disk is negligible \\citep{Martin1998}. Finally, RRLSs are pulsating stars with relatively short periods ($0.3-1.0$ d) and large light curve amplitudes ($V\\sim0.3-1.2$ mag), which makes them easily identifiable by a photometric multi-epoch survey \\citep[see e.g. \\citetalias{Vivas2004},][]{Kinemuchi2006}.   \\setcounter{table}{0} \\begin{table*} \\begin{minipage}{160mm} \\caption{Characteristics of recent large-scale RRLS surveys} \\begin{footnotesize} \\begin{tabular}{lcccccl} \\hline Survey & Filters & Area (\\sqdeg) & Completeness & Telescope & Observatory & Reference \\\\ \\hline QUEST  & $VRI$ & $476$ & $V<19$\\,\\,\\,\\, & $1.0$ m Schmidt & NOV & This Work \\\\ QUEST  (V04) & $V$ & $380$ & $V<19$\\,\\,\\,\\, & $1.0$ m Schmidt & NOV & \\citetalias{Vivas2004} \\\\ NSVS         & ROTSE-NT & $\\sim31000$       & $V<14$\\,\\,\\,\\, & $4\\times200$ mm ROTSE-I  & Los Alamos & \\citet{Kinemuchi2006} \\\\ ASAS-3      & $V$     & $\\sim31000$ & $V<14$\\,\\,\\,\\, & $2\\times200$ mm ASAS & Las Campanas & \\citet{Pojmanski2002} \\\\ LONEOS    & LONEOS-NT & $1430$        & $V<18$\\,\\,\\,\\, & $0.6$ m Schmidt & Lowell  & \\citet{Miceli2008} \\\\ SEKBO      & $B_MR_M$  & $1675$    & $V<19.5$ & $1.27$ m MACHO & Mount Stromlo & \\citet{Keller2008} \\\\ HATNET    & $I$ & $67$ & $V<15$\\,\\,\\,\\, & $11$\\,\\,\\,\\, cm HATNET & Fred L. Whipple & \\citet{Hartman2004} \\\\ OGLE        & $I$ & $0.87$ & $V<20.5$ & $1.3$ m Warsaw & Las Campanas & \\citet{Udalski1998} \\\\ SDSS Str82 & $ugriz$  & $249$     & $V<22$\\,\\,\\,\\, & $2.5$ m SDSS & Apache Point & \\citet{Watkins2009} \\\\ & & & & & & \\citet{Sesar2010}\\\\ MACHO     & $B_MR_M$  & $\\sim75$ & $V<20$\\,\\,\\,\\, & $1.27$ m MACHO & Mount Stromlo & \\citet{Kunder2008}, \\\\ & & & & & & \\citet{Alcock2003}\\\\ \\hline \\end{tabular} \\end{footnotesize} \\label{t:survey_details} \\end{minipage} \\end{table*}  RRLSs have been extensively used as tracers of the Galactic halo by numerous surveys which have studied their spatial distribution in a wide range of distances, from very near the Galactic centre up to large distances $R_{gal}\\ga100$ kpc. This is illustrated in Figure \\ref{f:survey_coverage_aitoff} which shows, in an Aitoff projection map in galactic coordinates, the footprint of recent large-scale RRLS surveys including the present one. Table \\ref{t:survey_details} summarizes the characteristics of each survey.  Unlike in the Galactic halo, the distribution of RRLSs in the thick disk has been less thouroughly studied, in particular due to the lack of deep RRLS surveys at low/intermediate galactic latitudes. As illustrated in Figure \\ref{f:survey_coverage_aitoff}, the Galactic Plane area $|b|<20\\degr$ has only been covered by the large-scale yet shallow ($V<14$ or $R_{hel}\\la 5$ kpc) ASAS-3 \\citep{Pojmanski2002} and NSVS \\citep{Kinemuchi2006} surveys, as well as the compilations of nearby RRLS ($V\\lsim13$) from \\citet{Layden1994,Layden1995} and \\citet{Maintz2005}; having been avoided by deep surveys ($16<V<20$). This makes ours, the first deep large-scale RRLS survey conducted at low Galactic latitudes. Furthermore, our survey probes the outer regions of the thick disk, outside the solar circle. In particular, \\citet{Layden1995,Maintz2005b} and \\citet[][NSVS]{Kinemuchi2006} have studied the distribution of thick disk RRLSs in the vertical direction, but the limited coverage at low galactic latitude prevented the determination of the scale length as well as the exploration of deviations from an exponential profile such as the \\emph{warp} and \\emph{flare}, which have been observed in the Galactic thin and thick disks with other tracers \\citep[e.g.][]{LopezCorredoira2002,Momany2006,Hammersley2011}. For these reasons it was essential to conduct a deep RRLS survey at low galactic latitude, which provided an ample spatial coverage both in the radial and vertical directions.  In this paper we will present the catalogue of the QUEST\\footnote{Acronym for the `QUasar Equatorial Survey Team' } RR Lyrae survey of the thick disk, which spans a total area of $476$ \\sqdeg at $-30\\degr<b<+25\\degr$, with multi-epoch $V,R,I$ observations obtained with the $1.0/1.5$m J\\\"urgen Stock  Schmidt telescope and the QUEST mosaic camera, at the National Astronomical Observatory of Venezuela (NOV). In Section \\ref{s:survey} we describe the observations and spatial and time coverage of the survey, as well as the photometric processing, error estimation and calibration. In Section \\ref{s:rrl_search} we describe the procedures used in the identification of RRLSs, we estimate the completeness and possible contamination of the survey and we describe the computation of periods, amplitudes, reddenings, photometric metallicities and distance for the identified RRLS. In Section \\ref{s:conclusions} we summarize the characteristics of our survey. The computation and analysis of RRLS density profiles for the Galactic thick disk and halo will be presented in an upcoming paper of the series. ",
        "conclusions": "\\label{s:conclusions}  We have presented the QUEST RR Lyrae survey at low galactic latitude, which spans an area of $476$ deg$^2$ on the sky, with multi-epoch observations in the $V$, $R$, and $I$ photometric bands for $6.5\\times10^6$ stars in the galactic latitude range $-30\\degr \\leqslant b \\leqslant +25\\degr$. The survey has a mean number of $30$ observations per object in each of the $V$ and $I$ bands and $\\sim25$ in the $R$ band, while in the best sampled areas each object can have up to $\\sim120-150$ epochs in $V$ and $I$ and up to $\\sim100$ in $R$. The completeness magnitudes of the survey are $V=18.5$ mag, $R=18.5$ mag and $I=18.0$ mag with saturations of $V=R=14.0$ mag and $I=13.5$ mag.  The survey has identified 211 RRLSs, 160 \\textit{bona fide} stars of type \\typeab~and 51 strong candidates of type~\\typec. The RRLS catalogue presented here contains the positions, mean magnitudes in $VRI$, periods, amplitudes, photometric metallicities, distances and individual extinctions computed from minimum light colours for each star, as well as the period-folded light curves for \\rrab~stars (Tables \\ref{t:rrab_lcparams}, \\ref{t:rrc_lcparams}, \\ref{t:rrab_physparams} and Figure \\ref{f:rrl_lcs}). The typical distance errors obtained were $\\sim7$ per cent. The completeness of the RRLS survey was estimated in $\\ga95$ per cent for \\rrab~ and $\\sim85$ per cent for \\rrc~stars, and was computed from light curve simulations reproducing the time sampling and other observing characteristics of our survey, as well as the variable identification and period searching techniques used.  Our RRLS survey spans \\emph{simultaneously} a large range of heliocentric distances $0.5 \\leqslant R_{hel}\\mathrm{(kpc)} \\leqslant 40$ and heights above the plane $-15\\leqslant z\\mathrm{(kpc)} \\leqslant +20$, with well characterized completeness across the survey area. Combining ours with a bright RRLS survey \\citep[such as  those from ][see Sec. \\ref{s:intro}]{Layden1995,Maintz2005} to increase the distance coverage even further, will result in a very good tool for studying the Galactic thick disk's structure, in particular for the determination of the density profile's scale length, scale height as well as their dependence with radial distance in order to explore the flare and the truncation profile of the thick disk. This way, the QUEST RRLS from this work will help put constraints on the structure of the more external parts of the thick disk toward the Anticenter, as well as the mechanisms that could have contributed to the formation of the Galactic thick disk. These issues will be addressed in detail in an upcoming paper of the series."
    },
    "1208/1208.6139_arXiv.txt": {
        "abstract": "We present a clustering analysis of near ultraviolet (NUV) - optical color selected luminosity bin samples of green valley galaxies. These galaxy samples are constructed by matching the Sloan Digital Sky Survey Data Release 7 with the latest Galaxy Evolution Explorer source catalog which provides $NUV$ photometry. We present cross-correlation function measurements and determine the halo occupation distribution of green valley galaxies using a new multiple tracer analysis technique.  We extend the halo-occupation formalism, which describes the relation between galaxies and halo mass in terms of the probability $P(N, M_{\\mr h})$ that a halo of given mass $M_{\\mr h}$ contains $N$ galaxies, to model the cross-correlation function between a galaxy sample of interest and multiple tracer populations simultaneously. This method can be applied to commonly used luminosity threshold samples as well as to color and luminosity bin selected galaxy samples, and improves the accuracy of clustering analyses for sparse galaxy populations.  We confirm the previously observed trend that red galaxies reside in more massive halos and are more likely to be satellite galaxies than average galaxies of similar luminosity. While the change in central galaxy host mass as a function of color is only weakly constrained, the satellite fraction and characteristic halo masses of green satellite galaxies are found to be intermediate between those of blue and red satellite galaxies. ",
        "introduction": "Most nearby galaxies fall into one of two well-known and well-characterized categories. They are either passively evolving elliptical galaxies with old stellar populations, red in color and typically living in high-density regions, or they are actively star-forming spiral galaxies with blue color. The latter often are field galaxies or reside in other low-density regions like cluster outskirts.  This blue/red galaxy color bimodality has been observed to be in place already around $z\\sim 1$. The fraction of red galaxies increases with time \\citep[e.g.,][]{F07} and therefore galaxies must transition from blue to red. Galaxies in this transitional stage characteristically show low levels of recent star formation. As ultraviolet emission is a sensitive tracer of recent star formation, these transition galaxies are easily identified in a $(NUV-r)$--$M_r$ color--magnitude diagram where they populate a ``green valley\" between well-localized red and blue sequences \\citep{W07}.  The relation between galaxy color and environment density also evolves with redshift, such that the fraction of red galaxies increases with time in dense environments but stays nearly constant for field galaxies \\citep[e.g.,][and references therein]{C07}. This indicates the transition from blue to red galaxies may be driven by environmental processes, associated with the infall of a galaxy into a larger halo (``cluster\"). Proposed mechanisms broadly fall into one of the following categories: galaxy--galaxy interactions, such as galaxy mergers, merger driven nuclear activity and high speed galaxy interactions, galaxy--intra cluster medium interactions (e.g., ram pressure stripping or thermal evaporation), and interactions between an infalling galaxy and the cluster potential (e.g., truncation through tidal forces). Observationally these are disentangled through their characteristic timescales, the dependence of their respective efficiencies on halo mass, and position within the cluster \\citep[][]{T03,C06, ME07}; for example, galaxy mergers are expected to be one of the dominant processes in group-scale halos and in the outskirts of massive clusters.  In the framework of $\\Lambda$ cold dark matter (CDM) cosmology, the evolution and spatial distribution of dark matter halos is relatively well understood. A common technique for inferring the masses of halos hosting different galaxy populations is to measure the angular or spatial clustering of galaxies and relate it to the predicted clustering and abundance of dark matter halos. While the relation between galaxy and dark matter clustering on large scales can be approximately described by scale-independent biasing, the situation is more complicated -- and more informative about the physical processes at work -- on small scales: At the level of individual halos, so-called halo-occupation distribution (HOD) models \\citep[e.g.,][]{BW02} describe the relation between galaxies and mass in terms of the probability that a halo of given mass contains $N$ galaxies of a given type. Then galaxy clustering, for example the two-point correlation function, is modeled as the sum of contributions from galaxy pairs residing in the same halo and from galaxy pairs living in different halos.  This method of interpreting galaxy correlation functions has been used extensively: For example, \\citet[][see references therein for previous/high-z studies]{Zehavi10} analyze the completed (DR7) Sloan Digital Sky Survey (SDSS), and find, in agreement with previous results, that at the amplitude of the correlation function increases with luminosity, and that at fixed luminosity redder galaxies are more strongly clustered, due to redder galaxies being satellites in more massive (and thus more biased) halos. Based on correlation function measurements over the redshift range $0.2<z<1.2$ from the Canada-France-Hawaii Telescope Legacy Survey, \\citet{Coupon12} also find red central galaxies to reside in more massive halos than average central galaxies in the same luminosity sample.  The clustering of $(NUV-r)$ color selected galaxies from the \\emph{Galaxy Evolution Explorer (GALEX)} survey has previously been studied by \\citet{H07}, who measure the angular correlation function; \\citet{H09} and \\citet{L10} analyze spatial clustering as a function of star formation history and color respectively. These authors find the clustering of green galaxies to have intermediate strength compared to blue and red galaxies and to have a scale dependence closer to that of red galaxies. At small scales their analysis is strongly limited by statistics due to the small number density of green valley galaxies, limiting their ability to constrain the 1-halo term.  We extend the HOD formalism to simultaneously model the cross-correlation functions (CCF) of a sparse luminosity bin galaxy sample with multiple more abundant galaxy populations to study the environment of local green valley galaxies. We consider luminosity bin samples of green valley galaxies as the physical mechanisms populating the green valley, i.e., quenching star formation in blue galaxies or rejuvenating red galaxies, may depend on halo mass and thus vary with galaxy luminosity. Compared to an autocorrelation function based clustering analysis, measuring the CCF between (sparse) \\emph{GALEX} selected galaxies and more abundant samples of SDSS galaxies reduces the shot noise contribution to our measurements, and also increases the effective volume probed beyond the combined \\emph{GALEX}-SDSS footprint\\footnote{We note that the increase in effective volume is limited to those regions of the SDSS footprint that are closer to the combined \\emph{GALEX}-SDSS footprint than the largest scales probed by the CCF. Due to the patchy geometry of the \\emph{GALEX} footprint, these regions cover nearly the entire SDSS footprint.}. Extending previous work on HOD models for CCFs \\citep[e.g.,][]{K10} to simultaneously fit the clustering of the galaxy sample of interest with respect to multiple tracer populations is particularly helpful for analyzing the clustering of luminosity bin samples, which are harder to constrain than the more frequently used luminosity threshold samples.This allows us to put the separate piece of information found by \\citet{H09} and \\citet{L10} into a coherent analysis including HOD modeling, and improve the statistics due to the larger survey area included in the newest data release.  Throughout this analysis we assume a flat $\\Lambda$CDM cosmology with $\\Omega_{\\mr m}  =0.3$ and $\\sigma_8  = 0.8$. Unless specified otherwise, all distances are coming and quoted in $\\mr{Mpc}/h$, and all absolute magnitude are given in $h=1$ units. ",
        "conclusions": "We introduced a new analysis and HOD modeling technique for galaxy cross-correlation functions using multiple tracer populations. This approach is particularly useful for interpreting the clustering of  sparse and/or luminosity bin selected galaxy samples of interest. It is advantageous for the analysis of sparse galaxy samples as considering the cross-correlation function with more abundant galaxy populations significantly reduces the statistical uncertainty.  While the galaxy number density provides strong constraints on the HOD of luminosity threshold samples, the HOD of luminosity bin samples is independent of the galaxy abundance; in this case considering the cross-correlation with multiple tracer populations is particularly useful as it provides an additional mass scale for the calibration of the luminosity bin HOD. An additional advantage of this method is that modeling the CCF between a color selected sample and a color independent sample does not require assumptions on the correlation between central and satellite colors.  This allows us to constrain the central galaxy HOD of color and luminosity bin selected samples for the first time. We apply this multiple tracer technique to analyze the clustering of $(NUV-r)$ color selected blue, red, and green valley galaxy samples. Our key result is that halo mass of central galaxies, satellite fraction, and halo mass of satellite galaxies increase with $(NUV-r)$ color at fixed luminosity.  While our results indicate that the clustering properties of green valley galaxies are consistent with them being an intermediate population between blue and red galaxies, the $(NUV-r)$ selected green valley galaxy samples in this analysis consist of only about one thousand galaxies and are too small to provide insight on the transition mechanism(s) at work. In particular, the HOD parameters which describe the abundance and distribution of color selected satellite galaxies, i.e. the slope of the satellite occupation function, $\\alpha_X$, and the color dependence of the radial satellite distribution, $c_X$, are poorly constrained by the data. With data from future galaxy redshift surveys, these parameters will provide information on the efficiency of star formation quenching as a function of halo mass and location within a halo. Furthermore, larger data sets will enable a detailed measurement of the redshift space correlation function and thus enable constraints on the infall stage and satellite orbits of transitional galaxies.  The reduced $\\chi^2$ values of the best-fit HODs in our analysis of color selected galaxy samples are relatively large, and our model is particularly insufficient to reproduce the clustering of galaxies in or near the Sloan Great Wall. Overall, it is not surprising that a five parameter HOD model does not fully describe the color dependent clustering of galaxies. While the HOD formalism works well to describe the overall relation between (color independent) galaxies and their halos, it is questionable if the strong assumptions implicit in the HOD formalism, such as the one-to-one relation between halo mass an bias, hold for each sub-population. Additionally, numerical and observational results indicate that the influence of massive halos may extend beyond $R_{200}$, e.g., through highly eccentric satellite orbits \\citep{Benson05, Wetzel11} and infall related shocks extending beyond the virial radius \\citep[e.g.,][]{Balogh00}, which is not easily incorporated in halo models.  Finally we note that \\citet{Behroozi10, Leauthaud11} recently proposed an improved HOD parameterization based on a detailed model for the relation between stellar mass and halo mass. Their results (figure 3 in \\citet{Leauthaud11}) indicate that halo masses derived from the HOD parameterization for luminosity threshold samples adopted in our analysis (equation (\\ref{eq:NM})) may be biased by up to 40\\%, with the main source of this discrepancy being the assumptions of a power-law form and constant scatter for the luminosity-halo mass relation. For luminosity bin samples, however, these assumptions are better justified, and we expect only small discrepancies between different HOD parameterizations."
    },
    "1208/1208.3000_arXiv.txt": {
        "abstract": "We introduce a new method for constraining the redshift distribution of a set of galaxies, using weak gravitational lensing shear. Instead of using observed shears and redshifts to constrain cosmological parameters, we ask how well the shears around clusters can constrain the redshifts, assuming fixed cosmological parameters. This provides a check on photometric redshifts, independent of source spectral energy distribution properties and therefore free of confounding factors such as misidentification of spectral breaks. We find that $\\sim 40$ massive ($\\sigma_v=1200$ km s$^{-1}$) cluster lenses are sufficient to determine the fraction of sources in each of six coarse redshift bins to $\\sim$11\\%, given weak (20\\%) priors on the masses of the highest-redshift lenses, tight (5\\%) priors on the masses of the lowest-redshift lenses, and only modest (20-50\\%) priors on calibration and evolution effects.  Additional massive lenses drive down uncertainties as $N_{\\rm lens}^{-{1\\over 2}}$, but the improvement slows as one is forced to use lenses further down the mass function.  Future large surveys contain enough clusters to reach 1\\% precision in the bin fractions if the tight lens mass priors can be maintained for large samples of lenses.  In practice this will be difficult to achieve, but the method may be valuable as a complement to other more precise methods because it is based on different physics and therefore has different systematic errors. ",
        "introduction": "Photometric redshifts are of key importance to current and future galaxy surveys.  In a typical application, galaxies might be binned according to photometric redshift and then some analysis conducted on these binned source sets.  Avoiding systematic error in this case requires knowledge of the true redshift distribution of each photometric redshift bin.  Ideally, photometric redshift methods themselves produce reliable confidence intervals for this purpose, but external verification is important for controlling and quantifying systematic errors.  For deep surveys, direct verification with a spectroscopic sample representative of the photometric sample is very difficult, and this has led to the development of other methods such as cross-correlating each photometric redshift bin with other photometric redshift bins (Schneider \\etal\\ 2006, Erben \\etal\\ 2009) or with spectroscopic redshift bins (Newman 2008, Matthews \\& Newman 2010).  Here we present a new and independent method for reconstructing source redshift distributions, using the shear from weak gravitational lensing.  The shear induced by a given lens grows with source redshift in a well-understoood way which depends on cosmographic parameters.  Other authors (Jain \\& Taylor 2003, Bernstein \\& Jain 2004) have suggested using this dependence to constrain the cosmographic parameters from the observed shear-vs-redshift relation around identified lenses. (Note that this is distinct from cosmic shear, for which the redshift dependence also involves the growth of structure.)  In this paper we reverse the question and ask how well the redshift distribution of a set of galaxies can be constrained by shear measurements around lenses at a range of redshifts, if the background cosmology is already well determined (see Appendix~\\ref{sec-cosmo} for a demonstration that cosmological uncertainties are negligible here).  This may prove useful for applications which take the background cosmology as given, and it also provides an internal consistency check for surveys.  Although the shear around any given lens is rather weak and might be expected to provide very little constraint, a very large survey such as LSST (Ivezi\\^{c} {\\it et al.} 2008) will contain billions of sources lensed by tens of thousands of galaxy clusters.  Assuming the photometry is uniform over the sky, the entire dataset can be used to constrain the source redshift distribution.  In this paper we use Fisher matrix formalism to estimate an upper bound on the effectiveness of this method, and we discuss challenges to implementing the method on real data.  Because this method requires a very large survey, we refer to LSST throughout.  We refer the reader to the LSST Science Book (LSST Science Collaborations {\\it et al.} 2009) for more details on LSST survey parameters. ",
        "conclusions": "}  We have outlined a new method which uses gravitational lensing by clusters and knowledge of the cosmology to constrain source redshift distributions.  Although this is an unorthodox use of the lensing-redshift information, it has some strengths and some relevant applications.  The method's strengths are: \\begin{itemize} \\item it is independent of other methods.  It relies on no assumptions about galaxy spectral properties or biasing.  \\item it is applicable to any set of galaxies.  One can determine the redshift distribution of a set of galaxies chosen in any number of ways, such as photometric redshift cut, pure color selection, or even X-ray or radio selection.  Multiwavelength observations are not necessary as they are for photometric redshifts.  \\item in the context of a large imaging survey, it requires little additional data.  Where lens spectroscopic redshifts are not available, photometric lens redshifts will do, and optical richness provides a sufficient lens-mass prior for the high-redshift lenses. However, pinning down the low-redshift end of the source distribution will require strong mass priors on low-redshift lenses.  \\end{itemize}  Our Fisher matrix forecast involves some approximations and caveats: \\begin{itemize}  \\item Lens morphology: we have assumed axisymmetric isothermal lenses. Although we have argued that other profiles would not yield substantially different forecasts, a precise experiment would include some sort of marginalization over profiles.  This marginalization might include allowances for non-axisymmetry.  \\item Foreground/background structure: we have assumed that shear from foreground/background structure is negligible in the ``footprint'' of each lens, and totally dominant outside.  In reality, unmodeled shear will be present in the footprint as well, and contribute to slightly looser constraints.  However, unmodeled shear is not an irreducible source of noise.  One can choose to model previously unmodeled structure and thus add information to the Fisher matrix. This entails a cost-benefit analysis rather than a hard limit.  \\item Magnification: the observed source population behind lenses will be slightly different than in other locations, due to lensing magnification.  If unmodeled, this may result in a nontrivial systematic error. However, because the lenses are being modeled, the model magnification at any given point provides an easily available basis for correction.  Deep fields which establish the source counts fainter than the main survey flux limit will be very useful for this purpose.  As with the foreground/background structure, it may be possible to add information to the Fisher matrix with improved modeling of the magnification.  \\item Cluster contamination: the areas around clusters differ from the general area of a survey in another respect, the presence of cluster members.  Care must be taken so that the source set whose shear is being measured is not contaminated by these members.  \\item Cosmology dependence of $\\beta$: Appendix~\\ref{sec-cosmo} shows that marginalizing over uncertainty in cosmographic parameters will not degrade the results presented here.  But using the cosmology as input here implies that the redshift estimates from this method should not feed in to some other method which estimates the cosmology unless the covariances are properly treated.  \\item Wide source redshift bins: given a lens redshift, $\\beta$ varies over the width of a source redshift bin.  Accurately modeling this effect will require some assumptions about the source redshift distribution inside the bins.  It may be possible to find a parameterization which reduces this difficulty.  \\end{itemize}  The source redshift distributions obtained with this method will be most useful in applications where the background cosmology would have been assumed anyway, such as galaxy evolution, and especially in applications which are very sensitive to photometric redshift outliers.  For example, measurements of the bright end of the galaxy luminosity function at high photometric redshift are extremely sensitive to catastrophic outliers, as a small fraction of low-redshift galaxies mistakenly put at high redshift results in many more ``high-luminosity'' galaxies at that redshift.  Surveys can use this method to independently derive the full redshift distribution of the galaxies in each photometric redshift bin, thus characterizing the number and distribution of outliers.  Indeed, the use of multiwavelength photometry to infer both source redshift {\\it and} source spectral/photometric properties can be problematic, and this method offers a possible workaround.  Other methods such as cross-correlation with a spectroscopic sample (Newman 2008, Matthews \\& Newman 2010) should be able to achieve 1\\% accuracy without assuming a cosmology or requiring expensive lens mass priors.  This method may therefore not be competitive statistically, but as it is based on different physics it may offer value in having different systematic errors.  Redshift distributions derived using this method can and should be used in consistency checks.  Given the best-fit cosmology from weak lensing tomography, this method should produce source redshift distributions which are consistent with those used in the weak lensing tomography.  A prime application for this method may therefore be as an internal consistency check for surveys."
    },
    "1208/1208.1834_arXiv.txt": {
        "abstract": "Different diagnostics of hot star wind mass-loss rates provide results that are difficult to reconcile with each other. The widely accepted presence of clumping in hot star winds implies a significant reduction of observational mass-loss rate estimates from diagnostics that depend on the square of the density. Moreover, the ultraviolet \\ion{P}{v} resonance lines indicate a possible need for even stronger reduction of hot star mass-loss rates, provided that \\ion{P}{v} is a dominant ionization stage of phosphorus at least in some hot stars. The latter assumption is challenged by a possible presence of the XUV radiation.  Here we study the influence of the XUV radiation on the \\ion{P}{v} ionization fraction in the hot star winds. By a detailed solution of the hydrodynamical, radiative transfer, and statistical equilibrium equations we confirm that sufficiently strong XUV radiation source may decrease the \\ion{P}{v} ionization fraction, possibly depreciating the \\ion{P}{v} lines as a reliable mass-loss rate indicator. On the other hand, the XUV radiation influences also the ionization fraction of heavier ions that drive the wind, leading to a decrease of the wind terminal velocity. Consequently, we conclude that the XUV radiation alone can not bring theory and observations in accord.  We fit our predicted wind mass-loss rates by a suitable formula and compare the results with the observational mass-loss rate diagnostics. We show that for supergiants and giants the theoretical predictions do not contradict the mass-loss rate estimates based on X-ray line profiles or density squared diagnostics. On the other hand, for main-sequence stars the predicted mass-loss rates are still significantly higher than that inferred from \\ion{P}{v} or X-ray lines. This indicates that the \"weak wind problem\" recently detected in low-luminosity main-sequence stars may occur to some extent also for the stars with higher luminosity. ",
        "introduction": "Mass-loss plays an important role in the massive star evolution. Most of the mass of massive stars is lost during their evolution from zero-age main sequence to the final remnant. Different processes contribute to the mass loss in individual evolutionary phases. These processes include line-driven winds during hot evolutionary stages \\citep{pulvina}, decretion disks in fast-rotating stars \\citep{los91,sapporo}, LBV-type of explosions in hot supergiants \\citep{vodovar}, dust-driven winds in cool supergiants \\citep{wojta}, and a final supernova explosion \\citep{medaci}. Consequently, estimates of amount of mass lost per unit of time (mass-loss rate) as a functions of stellar parameters belong to one of the most important ingredients of evolutionary models.  Unfortunately, the uncertainties in modern mass-loss rate determinations significantly affect the evolutionary models of massive stars. In the case of the line-driven wind of hot stars, these uncertainties seem to be mostly connected with the occurrence of small-scale inhomogeneities \\citep{potclump}. The inhomogeneities are typically divided into three differents groups (microclumping, porosity, and vorosity) according to their influence on the spectral features, although they may be caused by the same structure observed in different wavelengths. Microclumping (also frequently referred to as clumping), that accounts for the enhanced density in the optically thin inhomogeneities, can be most easily incorporated in the wind models \\citep[e.g.,][]{hagr,pulchuch,sanya}. Microclumping affects the ionization equilibrium via enhanced recombination, consequently it influences the radiative transfer only indirectly \\citep{abbob}. On the other hand, porosity (also referred to as macroclumping) accounts for the nonnegligible optical depth of inhomogeneities (which may become optically thick), and directly infuences the radiative transfer \\citep{lidaarchiv,sund,clres1}. Vorosity affects the line profiles, and it is connected with the different Doppler shifts of individual inhomogeneities \\citep{ocvor}. From the point of view of mass-loss rate predictions, the inhomogeneities may affect the ionization fractions of wind driving ions \\citep{sanya,muij}, or they may lead to the decrease of the wind mass-loss rate due to the base turbulence \\citep{lucyinthewind,cmf1}.  There could be a simple observational solution of the mass-loss rate determination problem: to find such observational characteristic that is not affected by the wind inhomogeneities. There are two potential candidates for such convenient observables: the X-ray radiation \\citep[either line profiles or the continuum flux distribution,] []{mcf,owco,tlustocerv,coh}, and unsaturated resonance line profiles. The latter case is fulfilled for trace elements, and especially interesting is the \\ion{P}{v} ion \\citep{prvnifosfor,maso,full}. However, even these characteristics face some problems. The opacity in the X-ray domain scales mostly linearly with the density, consequently its effect on the X-ray diagnostics should be in principle possible to model in a straightforward way. Indeed, the observations of mostly symmetric X-ray line profiles, that are not strongly affected by absorption, indicate low wind mass-loss rates \\citep{celakoza,cohzet,gagne}. These results are, however, challenged by a possible effect of the  porosity on the X-ray opacity \\citep{lidarikala,oski}. Note that there is not a general consensus on this problem \\citep{taulida,tauowo}.  Here we concentrate mostly on the other promising observational characteristic, i.e., on the \\ion{P}{v} resonance lines. The weakness of observed \\ion{P}{v} lines also indicates low mass-loss rates of hot star winds \\citep{full}. However, the ionization fraction of \\ion{P}{v} may be modified by the effect of clumping \\citep{pulamko,sanya}. Additional changes of the \\ion{P}{v} ionization fraction may be caused by the influence of X-rays. Although our previous calculations indicated that \\ion{P}{v} ionization fraction is not strongly affected by X-rays \\citep{nlteiii}, \\citet{walcp} argued that the extreme ultraviolet radiation (hereafter XUV) may affect \\ion{P}{v} ionization fractions. Since the work of \\cite{walcp} was based on rather simplified ionization estimates and the calculations presented in \\citet{nlteiii} were done without enhanced XUV radiation, we decided to fill this gap and we apply our NLTE wind models to study the influence of the XUV radiation (parameterized in a convenient way) on the \\ion{P}{v} ionization fraction. The XUV region is defined here as the energy from interval 54.4\\,eV (\\ion{He}{ii} edge) to 124\\,eV. ",
        "conclusions": "We tested the influence of XUV radiation on \\ion{P}{v} ionization fractions in hot star winds. Using hydrodynamical NLTE wind modelling we confirmed the conclusion of \\citet{walcp} that the XUV radiation may decrease the \\ion{P}{v} ionization fraction. As the result of various ionization and recombination processes, the \\ion{P}{v} ionization fraction in hot star winds does not come close to the value of 1, implying that the role of the \\ion{P}{v} lines as mass-loss rate indicators in hot star winds is not so straightforward. On the other hand, the large amount of XUV radiation necessary to significantly lower the \\ion{P}{v} ionization fraction to the values implied by observations leads to a decrease of the wind terminal velocity due to inefficient line driving. This contradicts the observations. Moreover, the wind clumping has an opposite effect than the XUV emission decreasing the wind ionization. Consequently, it is unlikely that the effect of lowering the \\ion{P}{v} ionization fraction by XUV radiation alone can bring the theory and observations in accord.  We also provide a useful mass-loss rate formula and compare it with other mass-loss rate diagnostics. We show that for supergiants and giants this formula passes two important observational tests against mass-loss rates derived from X-ray line profiles and $\\rho^2$ diagnostics. This supports the reliability of mass-loss rates predictions derived from modern wind codes for luminous hot stars.  For main-sequence stars the predicted mass-loss rates are significantly higher than those inferred from the \\ion{P}{v} or X-ray lines. This may indicate that the \"weak wind problem\" recently detected in low-luminosity main-sequence stars occurs to some extent also for the stars with higher luminosity. While the explanation of this \"weak wind problem\" for low-luminosity main-sequence stars is at hand (too large cooling length, \\citealt{luciebila}, \\citealt{martclump}, \\citealt{cobecru}, \\citealt{nlteiii}, \\citealt{lucyjakomy}, or influence of X-rays on the mass-loss rate, \\citealt{nemajpravdu}) such explanation valid for all main-sequence stars with any luminosity is currently missing."
    },
    "1208/1208.0003_arXiv.txt": {
        "abstract": "We study the effects of radial flows on Galactic chemical evolution. A simple analytic scheme is developed prescribing the coupling of infall from the intergalactic medium and radial flows within the disc based on angular momentum conservation. We show that model parameters are tightly constrained by the observed \\feh-abundance gradient in the Galactic disc. By this comparison the average rotational velocity of the onfalling material can be constrained to $ 0.7 \\le v/\\Vc \\le 0.75$, or respectively $\\sim 160 \\eunitfrac{km}{s}$ when assuming a constant disc circular velocity of $\\Vc = 220 \\eunitfrac{km}{s}$. We test the robustness of this value against the influence of other processes. For a very simple model of inside-out formation this value changes only by $\\Delta v/\\Vc \\sim 0.1$, i.e. $\\sim 20 \\eunitfrac{km}{s}$, and significantly less on more realistic scenarios, showing that inside-out formation does not alone explain the abundance gradient. Effects of other uncertain parameters, e.g. star formation history and star formation efficiency have very small impact.  Other drivers of inflow beyond our explicit modelling are assessed by adding a fixed inflow across the whole disc. The churning amplitude only mildly affects the results mostly by slightly flattening the metallicity gradient in the inner disc. A new process causing radial gas flows due to the ejection of material by stars moving on non-circular orbits is studied and seems to contribute negligibly to the total flows.  We further show that gaseous outer discs cannot be the main source feeding the persistent star formation in the inner regions by a direct inflow. ",
        "introduction": "Radial flows of gas have been studied extensively in models of chemical evolution. Their importance has first been recognised by \\cite{TinsleyLarson1978}, who emphasise their effects on galactic abundance gradients and conclude that these cannot be ignored in any viable model. \\cite{LaceyFall1985} discussed three main drivers of these flows. First, onfalling gas might have a different specific angular momentum than the gas rotating in the disc. As the onfalling gas is suspected to have a lower rotational velocity, these flows are expected to be directed inwards at a velocity in the range of a few\\,\\eunitfrac{km}{s} \\citep{LaceyFall1985}. Second, friction within the rotating gas layers might cause radial flows. A viscous disc is subject to an angular momentum flux whenever the angular speed changes with radius. A classic Galactic disc with its nearly flat rotation curve has a global outwards directed angular momentum flux, which drives inflows in the inner regions and an expansion of its outskirts. Estimates of this effect yield velocities of the order of 0.1\\eunitfrac{km}{s} \\citep{LaceyFall1985}. Third, interactions of the gas with a central bar and spiral patterns, especially around the corotation resonance, will cause radial flows, typically directed away from the position of the resonance. \\citet{LaceyFall1985} cite velocities of 0.3\\eunitfrac{km}{s} at a distance of 10\\eunit{kpc} from the galactic centre. Further the gas loses angular momentum to the stellar population by interactions/shocks along the spiral arms. A comparison of these quantities shows that especially stronger flows are preferentially driven by angular momentum dilution via onfall.  Infall has been known to affect the evolution of disc galaxies for many years \\citep{Larson1972}. Even if no gas at all was lost to the intergalactic medium, roughly one third of the mass involved in each star formation cycle would get locked up in stellar remnants or very long-lived low mass stars and the star formation in spiral galaxies would exhaust the available gas reservoirs on cosmologically short timescales \\citep{Larson1980}. Hence the observed star formation rates require ample infall of fresh gas. Infall is also required from a chemical evolution point of view to avoid exceedingly high metallicities after a few Gyrs and to avoid deuterium abundances far below the observations \\citep{Linsky2006}. The mode of this accretion is obscure: It is frequently adopted that gas is delivered directly to galaxies by infalling satellites. This direct infall is observed \\citep{Sancisi2008}, but they also find that it brings about an order of magnitude too little material to sustain the current star formation rates. Further, such ``cosmological infall'' leads to rather unsatisfacory outcomes \\citep[see][]{Colavitti08} in chemical evolution models, since the clumpy accretion and star formation produce unobserved excursions in the detailed abundances, especially in the $\\feh-\\afe$ plane.  Analytical chemical evolution models assumed a vanishing angular velocity of onfalling material \\citep{MayorVigroux1981}. Later onfalling material with a constant fraction of the disc's angular velocity at the given radius was considered as well \\citep{PittsTayler1989,PittsTayler1996}. But no argument aside from trying to fit the observed abundance gradients could be made for any particular prescription. In a different line of development radial gas flows were implemented to classical chemical evolution models, starting with \\cite{GoetzKoeppen}, later \\cite{Portinari00} and recently by \\cite{Spitoni11}. Those models assume preset patterns of the inflow velocity in the disc, but do not draw a significant link to physical properties like the angular momentum budget.  It seems natural, though, to link the available angular momentum budget to the puzzle of Galactic accretion. The deficit of visible cold gas infall was recently connected to the presence of a concurrent mode of delivery: Hot coronal gas can be swept up by turbulent cooler gas clouds that are driven out of the Galactic plane by supernova feedback \\citep{FB06,FB08,MB10}. Dynamic arguments and observations strongly suggest that the coronal gas should lag behind the rotation speed of the disc \\citep{Marinacci2011}. Theoretical calculations for the Milky Way predict values for the global accretion rate sufficient to sustain its star formation and put an estimate of the lag of accreted coronal material at around $75 \\eunitfrac{km}{s}$ compared to the disc rotational velocity of $220 \\eunitfrac{km}{s} $ \\citep{Marasco2011}. Hence in contrast to cosmological cloud accretion this offers a firm framework with relatively well constrained properties of accretion. In addition they derive a specific profile for the accretion of gas onto the Galactic disc.  A process neglected so far is the systematic lag of stellar yields themselves behind the local gas velocities as well as friction between the stars and the interstellar medium. This is mediated by the larger velocity dispersion and hence asymmetric drift of stars. We will discuss this in the following paper.  The purpose of this paper is to give a better parametrisation of flows than that used originally in the model of SB09 in the sense that it provides sufficient freedom to fit the data and depends on physically meaningful parameters. The model crucially relies on the assumption of conservation of angular momentum which may not be true in all cases, but which should be a plausible first approximation to understand the overall flows in the Galaxy. Our approach is different from those found in the literature as we assume a radially varying profile for the angular velocity of the onfalling gas. Fitting the model to the data provides information on the kinematic properties of the onfalling gas. These properties depend on the proposed origin of the onfalling material and the assumed way of accretion onto the Galaxy. A comparison of the predictions derived from a model of chemical evolution with those of theories for the origin of the onfalling material might help to shed light on the question of how galaxies accrete gas.  The paper is organised as follows. Section \\ref{sec:SB09} summarises the aspects of the SB09 model \\citep{SBI} important to this work. In section \\ref{sec:modeleq} we give the equations of the new flow model to be used in the simulations. We model the angular velocity of the onfalling gas as a linear function of galactocentric radius. The general behaviour of the model and the dependence on its parameters is presented in section \\ref{sec:modelbe}. In section \\ref{sec:Results} we fit the new model to the Galactic radial abundance gradient to constrain the physical model parameters. We investigate the implications of the accretion profile proposed by \\cite{Marasco2011} in section \\ref{sec:accprofile}. In section \\ref{sec:dyingstars} we study the proposed new mechanism for causing radial flows due to the lag of stellar yields. Section \\ref{sec:inside-out} discusses an implementation of inside-out formation in the model and section \\ref{sec:origin} explores limits on the accretion from an extended gas disc. Section \\ref{sec:sum} sums up. ",
        "conclusions": "\\label{sec:sum}  While recent works concentrated mostly on describing the effects of inflows of a certain preset velocity, we coupled the inflow process to conservation of angular momentum between the accreted material and the gas already present in the disc. In this picture, the angular momentum of the disc gas is diluted by freshly accreted material with lower specific angular momentum driving a radial flow through the disc. This coupling is used to derive the spatial structure of the radial flows consistently within the model. The resulting flow and onfall pattern is not known from the beginning, but rather a prediction of the model. Due to its physical motivation the new prescription is preferable to the older prescription of SB09 and also to models with a preset inflow.  We have shown that the chosen prescription can produce a vast range of radial metallicity distributions in disc galaxies and studied the dependence of the observables on different assumptions about the qualities of the inflow.  We fitted the model to observations of the present day metallicity gradient of our Galaxy by \\citet{Luck2011II}. This comparison constrained the average rotational velocity of the onfalling material to be $ 0.7 \\le v/\\Vc \\le 0.75$ corresponding to $154\\eunitfrac{km}{s} \\le v_{onfall} \\le 165 \\eunitfrac{km}{s}$ assuming a constant $\\Vc = 220 \\eunitfrac{km}{s}$. This value compares well with the predictions for an accretion of gas from the corona (when neglecting vertical velocity gradients) at about $V_{corona} \\approx 145\\eunitfrac{km}{s}$ \\citep{Marinacci2011, Marasco2011}. However, the specific accretion profile predicted by \\cite{Marasco2011} is not compatible with our simulations, since the rather sharp peak in Galactocentric radius would drive flows that create a radial metallicity distribution at odds with the data, if not some other process fills the gap especially in the inner disc regions.  One major uncertainty is how much angular momentum is given up by the gas itself through various channels. We studied a so far neglected process where the ejection of material from stars moving on non-circular orbits dilutes the angular momentum via their asymmetric drift. If there is not an unknown amplification, this momentum exchange has only a very small impact on radial flow velocities of the order of $v_{flow} \\sim 0.01\\eunitfrac{km}{s}$ and can be safely neglected. More of a concern are the well-known processes of gravitational interaction of the gas with spiral arms or the bar and possibly friction within the gas layers. Considering the estimates by \\cite{LaceyFall1985} they might play a role in additionally changing the gradient. We have considered models in which a fixed constant flow velocity was added to our flow model to describe part of the uncertainty of these processes. We conclude that flows at that level certainly change our results quantitatively if they extend across the whole Disc and persist for the whole age of the Galaxy, but for values within the estimates the total effects remain small of order $0.05 \\, v/\\Vc$. However, the local structure and gas dynamics may well be affected considerably without affecting the global flow levels and the gradient.  Resonances stable in galactocentric radius are another source of uncertainty to the flow pattern. Any significant influence, however, should show up as an excursion in the radial abundance profile and up to now is unobserved, making such a bias a minor concern.  Radial migration \\citep{SellwoodB01} has only a minor influence in flattening the inner disc abundance gradient slightly via mixing. Even larger amplitudes are easily compensated by the inflow.  These results were shown not to depend strongly on the resolution or the churning amplitude. Additional dependence on model parameters, specifically the assumed infall history and the star formation efficiency, was found to be relatively small.  Declining star formation rates tend to reduce the gradient in our models. As the radial flows are directly coupled to the accreted material which determines the star formation rates, a declining SFR also means decreasing flows. For the specific model parameters, a very fast decline of the SFR shifts the gradient to lower values, but a slower decline has relatively little effect in steepening the gradient.  The star formation efficiency mostly affects the simulation by controlling the total amount of gas in the disc. For different levels of radial flows, the effects on the gradient are qualitatively different, but for the region corresponding to the observed data a higher gas mass tends to flatten the gradient, hence needing stronger flows and a lower angular momentum of onfalling material to give the correct value.  We also examined the impact of inside-out formation. Generally inside-out formation raises the ratio of metal production to infall of fresh material in the inner disc relative to the outer disc facilitating a steeper metallicity gradient. However, a simple linear growth of the gas scale length over the complete age of the Galaxy is strongly limited by observations, and a model in which the gas disc mainly grows in the early times of the simulations is more adequate in order to reproduce the stellar disc scale length measured today. In this latter case the effects on the present-day metallicity gradient are greatly reduced. It was impossible for us to achieve the measured abundance gradient just by inside-out formation and so conclude that radial flows are essential for explaining the observations.  Independent from these considerations we showed that the large amounts of HI gas in extended layers beyond the stellar disc cannot feed the majority of fresh material required to the inner disc, since this would imply too large radial flows and a too steep abundance gradient."
    },
    "1208/1208.1763_arXiv.txt": {
        "abstract": "We use  a sample  of $z\\sim3$ Lyman  Break Galaxies (LBGs)  to examine close pair clustering statistics in comparison to LCDM-based models of structure formation.  Samples are  selected by matching the LBG number density,  $n_g$, and  by  matching the  observed  LBG 3-D  correlation function  of  LBGs  over  the  two-halo term  region.   We  show  that UV-luminosity abundance  matching cannot reproduce  the observed data, but if  subhalos are  chosen to reproduce  the observed  clustering of LBGs we are able to reproduce the observed LBG pair fraction, ($N_c$), defined as  the average number  of companions per galaxy.   This model suggests an over  abundance of LBGs by a factor  of $\\sim5$ over those observed, suggesting  that only  $1$ in $5$  halos above a  fixed mass hosts  a  galaxy  with  LBG-like  UV  luminosity  detectable  via  LBG selection  techniques.   This overdensity  is  in  agreement with  the results of a Millennium 2 analysis and with the discrepancies noted by previous  authors using  different types  of simulations.   We  find a total observable close  pair fraction of $23 \\pm  0.6$ per cent ($17.7 \\pm  0.5$)  per cent  using  a  prototypical  cylinder radius  in  our overdense fiducial model and $8.3 \\pm 0.5$ per cent ($5.6 \\pm 0.2$ per cent)  in an abundance  matched model  (impurity corrected).   For the matched spectroscopic slit analysis, we find $N_{cs}(R) = 4.3 \\pm 1.55 (1.0 \\pm 0.2 )$ per cent and  $N_{cs} = 5.1 \\pm 0.2$ ($1.68 \\pm 0.02$) per cent, the average number of companions observed serendipitously in randomly  aligned spectroscopic slits,  for fiducial  slits (abundance matched), whereas the observed fraction of serendipitous spectroscopic close pairs  is $4.7 \\pm 1.5$ per  cent using the full  LBG sample and $7.1 \\pm  2.3$ per  cent for a  subsample with  higher signal-to-noise ratio.  We conduct the same analysis on a sample of dark matter haloes from the Millennium  2 simulation and find similar  results.  From the results and an analysis of the observed LBG 2-D correlation functions, we show that the standard method of halo assignment fails to reproduce the break, or  up turn, in the LBG close pair  behavior at small scale ($\\lesssim 20 ~\\hkpc$ physical).   To reconcile these discrepancies we suggest that  a plausible fraction of  LBGs in close  pairs with lower mass (higher  density) than our  sample experience interaction-induced enhanced star  formation that boosts their  luminosity sufficiently to be  detected in  observational  sample  but are  not  included in  the abundance matched simulation sample. ",
        "introduction": "\\label{sec:intro}%  One of the fundamental predictions of a \\lcdm model of the universe is the hierarchical growth of structure.  However, direct observations of galaxy mergers,  and, by extension, statistics on  galaxy mergers, are difficult to obtain  due to the long time-scales  of the galaxy-galaxy merger process.  Correlations between galaxy characteristics and their environment suggest  that interactions play  a role in  setting galaxy properties   such   as  star   formation   rate,  colour,   morphology \\citep[e.g.][]{ToomreToomre72,   LarsonTinsley78,   Dressler80,  PG84, Barton:00,  Barton:03}.  However,  observational studies  of mergers and interactions can be difficult due to the low luminosities of tidal features and the difficulties  in quantifying galaxy morphologies.  At high  redshifts, $z  \\ge 1$,  these  problems are  exacerbated by  the decreased  apparent luminosity  and resolution  of the  galaxies being studied.  Since the studies of \\citet{Holmberg1937} close pairs of galaxies have provided an important  tool for the evaluation of  galaxy merger rates by providing counts  of merger candidates, and for  theories of galaxy formation  due to the  importance of  galaxy-galaxy mergers  in galaxy evolution.   Close galaxy  pair counts,  or counts  of morphologically disturbed systems, have  not only been used to  provide candidates for galaxy mergers,  but have  been used in  attempts to probe  the galaxy merger  rate   and  its  evolution   with  redshift  \\citep[][]{z&k89, Burkey94,  Carlberg94, Woods95,  Y&E95,  P97, Neus97,  Carlberg2000, LeFevre2000,  P2002,  Conselice03,  Bundy04, masjedi05,  Belletal06, Lotz06, Lin04}.  In \\citet{Berrieretal06} we present a method to analyse the close pair fraction of galaxies in a  simulation environment, with the close pair fraction ($N_c$) defined as the number of galaxies in close pairs in a volume  of space normalised  by the  total number  of galaxies  in the sample.  This  analysis demonstrates  the viability of  estimating the observable close pair fraction in simulations using simple criteria to assign galaxies to dark matter haloes.  \\citet{Berrieretal06}  argue that the  close luminous  companion count per galaxy does  not track the distinct dark  matter halo merger rate. Instead, it  tracks the luminous  galaxy merger rate.  While  a direct connection between the two has often been assumed, there is a mismatch because multiple galaxies  may occupy the same host  dark matter halo. The same  arguments apply  to morphological identifications  of merger remnants, which also  do not directly probe the  host dark halo merger rate.   This still leaves  close galaxy  pairs as  a tracer  of galaxy evolution and as a proxy of the galaxy merger rate.  At high redshift, the dense environment and smaller fraction of galaxy clusters   [where  the   large  velocity   dispersion   prevents  many satellite-satellite mergers, e.g. \\citet{Berrieretal2009}] mean that close pairs of galaxies are  likely to indicate actual mergers, though estimates of the timescales of these mergers may still be rather large \\citep[see   e.g.,][]{Kitzbichler2008,  Bertone09}.    As   a  result, observations of this process for  galaxies at high redshift are highly desirable  for  constraining the  high  redshift galaxy-galaxy  merger rate.  The Lyman break galaxies  (LBGs) are star forming galaxies efficiently identified using  colour selection criteria \\citep[e.g.,][]{Steidel96} and  comprise  a large  fraction  of  all  luminous galaxies  at  high redshift  \\citep[e.g.,][]{Reddy05,Marchesini2007}.   To  date,  a  few thousand LBG  spectra and tens of thousands  of photometric candidates have been obtained, making  LBGs a useful, well-studied population for high  redshift galaxy  spatial distribution  and close  pair analysis. \\citet{Conroy2006} used  subhalo abundance matching  techniques (SHAM) such  as those used  in \\citet{Berrieretal06}  and here,  to calculate angular  correlation functions  for LBGs  at high  redshifts, $z=3,4$. This work suggests  that abundance matching techniques may  be used to sample LBG  populations and statistics in simulations.   SHAM has been tested in  a variety of situations  at both low and  high redshift and has  been shown  to  be a  reasonable  tool for  matching galaxies  to populations of  dark matter halos  and generating halo mass  - stellar mass  relations  \\citep{Conroy2006, Vale&Ostriker2006,  Berrieretal06, Stewart2009,   Simha2010,   Guo2010,   Moster2010}.   Because   this technique  has been  suggested, and  indeed used,  as a  probe  of LBG clustering  statistics, it  will provide  our starting  point  in this analysis.   We also  explore  matching dark  matter  halo and  subhalo correlation   functions  to  the   observed  clustering   of  $z\\sim3$ LBGs. This technique  is similar to the work  of \\citet{Conroy2008} on $z\\sim2$ star forming galaxies.  In  this  paper,  we  use  a numerical  $N-$body  simulation  with  an analytically   generated  substructure,   adopting  the   approach  of \\citet{Berrieretal06}, to  compare close companion  counts directly to the observed  companion count for our  sample of $z \\sim  3$ LBGs from the survey  of \\citet[][hereafter S03]{Steidel2003} and  the survey of \\citet[][hereafter C05]{Cooke05}.   Our purpose is to  test the simple and   popular  \\citep{Conroy2008,Stewart2009,Simha2010}   theory  that galaxies  live   in  subhalos   and  that  UV   luminosity  correlates monotonically with halo mass/maximum  circular velocity at the time of accretion.  The structure  of this paper is as  follows.  In \\S~\\ref{sec:methods}, we  outline   our  methods,   discuss  our  observational   sample  in \\S~\\ref{sec:observations},          our         simulations         in \\S~\\ref{sec:simulations},  the  models  used  for  the  assignment  of galaxies to haloes in \\S~\\ref{sec:model}. The definitions of the close companion fraction, the photometric companion fraction, and the sample impurity  and   number  density  are  covered   in  \\S~\\ref{sec:nc}  - \\S~\\ref{sub:NG}  respectively.   We present  our  predictions for  the companion fraction, $\\Nc$, in  \\S~\\ref{sec:results}.  We begin with an examination of  $N_c$ from  $z=0-3$ with an  emphasis on  a comparison between  our simulations  and  the observational  values  at $z=3$  in \\S~\\ref{sub:evolution}.  The angular photometric close companion count is  the  topic  of  \\S~\\ref{sub:pred}.   Comparisons  with  previously existing close companion counts  are made in \\S~\\ref{sub:compare}.  We return to the number  density issue in \\S~\\ref{sub:Density}.  Finally, we discuss the implications of our results in \\S~\\ref{sub:Discussion}. We conclude with a summary in \\S~\\ref{sec:conclusions}.  In this  work we assume a  flat universe with a  standard cosmology of $\\Omega_m= 1 - \\Omega_{\\Lambda} = 0.3$, $h=1.0$, and $\\sigma_8 = 0.9$. ",
        "conclusions": "\\label{sec:conclusions}%  We have matched the 3-D two  point correlation function of a sample of $z   \\sim  3$  LBGs   to  a   sample  of   haloes  from   our  primary numerical/analytic cosmological simulation.  Using this sample we have mocked   observations  of   simulated  spectroscopic   slits   and  of photometric observations  of these galaxies.   We also test  our model with  data from the  Millennium-2 simulation  for verification  of our results.  This work has led us to several interesting results which we summarise in the points below, see also Table \\ref{table1} above.  \\begin{itemize}  \\item We demonstrate that neither standard SubHalo Abundance Matching (SHAM) or a two-point correlation function and mass matching scheme completely reproduces the observational  results. Neither  model can reproduce   galaxy   clustering     features  and   $n_g$   at  same time. Explicitly, we find that the standard  SHAM does not reproduce the  serendipitously  observed $N_{cs}$, and  the  break  in the LBG correlation  function  at  very   small   scales ($\\lesssim20~\\hkpc$ physical, $\\sim80~\\hkpc$ comoving).  \\item The number  density of our candidate LBG  sample is $\\sim$$4.75$ times the observed LBG number density.  The implication is that only $\\frac{1}{5}$ halos  above a fixed mass  are detectable LBGs.  These results    are     consistent    with       the     results       of \\citet{Nagamine2004,Nagamine06,Lee08, Dave2000,Lacey2011} and others which find a similar overdensity using other types of simulations.  \\item We find an observed close pair fraction $N_c = 0.228 \\pm 0.006$, which implies  an impurity corrected  close pair fraction of  $N_c = 0.177 \\pm 0.005$, ($\\sim18$ per cent) This result is consistent with the   previous   results   of   \\citet[][]{Conselice03,   Bertone09, Bluck09,Law12,    Lotz06,    Forster-Schreiber09}   considering  the uncertainties specific  to those  studies and  pair  fraction/merger rate assumptions.  \\item Our  simulated matched spectroscopic slits produce  a close pair fraction of  $N_{cs} = 0.0506  \\pm 0.0017$ for our  fiducial sample, defined to have a maximum  velocity separation of $\\pm 500$ \\kms and an on the  sky separation of $\\le 30.0 ~\\hkpc$.   This is similar to the observed fraction of  serendipitous spectroscopic close pairs of $N_{cs} =  0.047 \\pm 0.015$ for  the full observed  LBG sample and $N_{cs} =  0.071 \\pm  0.023$ for the  highest signal to  noise ratio sample. After correcting for false close pairs which may be observed we find $N_c = 0.0419 \\pm 0.0015$ in the simulation.  \\item If we  examine our catalogs using randomly  selected slitlets to generate a more  generic result, we find that  the expected fraction of  LBG pairs  that fall  serendipitously  into the  slitlets to  be $N_{cs} = 0.0430 \\pm 0.015$ before correcting for sample impurity.  \\item We find  a photometric close pair fraction of  $N_p = 0.336 \\pm 0.009 $  and after correcting for  the sample  impurity we find $N_p \\sim 0.189  $.  The latter  fraction reflects the ``real'' number of close pairs and therefore the real number of potentially interacting galaxy pairs.  As  mentioned above, only a portion  of the corrected value will be observable LBGs.   The difference in $n_g$ is a factor of $\\sim 4.75$ leading to a  corrected value of $N_p \\sim 3.99 $ per cent, which is consistent with our photometric pair fractions $N_p = 0.047  \\pm0.035$   estimated  from  our survey  and   the  survey of \\citet{Steidel2003}.  \\item  The analysis  of the  sample taken  from  Millennium-2 produces similar results to our primary simulation.  In Millennium-2, we find the   correlation  function-matched  sample   to  be   overdense  in comparison with  the observations by  a factor of  $\\sim4.25$.  This selected  sample  produces  a  $N_c  = 0.1521  \\pm  0.0006$  in  the cylinder, and $N_{cs}  = 0.033 \\pm 0.001$ in  the slits.  The sample shows an impurity of $I =  0.1433 \\pm 0.0176$ in the cylinder and $I = 0.1167  \\pm 0.0202$  for the slit,  producing corrected  values of $N_c = 0.1303 \\pm 0.0059$ and  $N_{cs} = 0.02919 \\pm 0.0014$ for the full cylinder and the spectroscopic slit respectively.  For our line of sight mock photometric sample  we find $N_p = 0.2299 \\pm 0.0144$, without corrections for impurity and $N_p = 0.1370 \\pm 0.0154$ after correction.  \\end{itemize}  The  excess  of   close (interacting)   pairs $\\le    20$ h$^{-1}$ kpc (physical) and the inability for the  standard abundance matching with monotonic UV-mass  halo assignment to describe the  steep slope in the observed  LBG  correlation  function  at very   small  scales provides insight into triggered star   formation and the detectability  of LBGs (and  LAEs) at $z=3$.  Our results  imply that  the spectroscopic slit close   pair fraction  and the  break    in the correlation   function represent the detection  of either a fraction  of less massive (higher density) LBGs with luminosities below  our magnitude cut  ($m_R=25.5$) as  a result of  an  enhancement in  luminosity from interactions, the ``turn-on\" of massive  haloes with  previous low  star formation as  a result of interaction, or, likely, a combination of both cases.  We find that LBGs likely represent $\\sim20-25$ per cent of all massive ($V_{in} > 133$ km sec$^{-1}$) haloes at $z\\sim3$ based on the results of  the analysis  of    our simulation, the   Millennium-2 simulation, simulation  analyses by  several other  authors, and the  power of our simulation analysis to predict the  close pair fraction from $z=1$  to $z=0$.  The full census of  detected star forming galaxies selected by various criteria suggest that  LBGs likely account for $\\gtrsim50$ per cent of the  massive haloes at  $z\\sim3$.   The remaining  fraction is likely   populated by  systems  with  low star  formation rates and/or systems that   are  not detected using  current  selection techniques. DLAs are a  promising  means  to explore   the remaining fraction   of massive haloes because they probe  galaxy haloes randomly, independent of luminosity, they have a high number  density, and they are found to reside   in   massive  haloes  \\citep{Cooke06,  Fynbo2003,  Fynbo2008, Fynbo2010, Fynbo2011, Moller2002, Moller2004, Schaye2001}.  The statistics  generated from our  mock spectroscopic slits  with the serendipitously  confirmed close  pairs from  observations  provides a potentially powerful tool to estimate the behaviour and nature of LBGs and the enhanced star formation rate from LBG interactions."
    },
    "1208/1208.1280_arXiv.txt": {
        "abstract": "We are systematically surveying all known and suspected Z Cam-type dwarf novae for classical nova shells. This survey is motivated by the discovery of the largest known classical nova shell, which surrounds the archetypal dwarf nova Z Camelopardalis. The Z Cam shell demonstrates that at least some dwarf novae must have undergone classical nova eruptions in the past, and that at least some classical novae become dwarf novae long after their nova thermonuclear outbursts, in accord with the hibernation scenario of cataclysmic binaries.  Here we report the detection of a fragmented \"shell\", 3 arcmin in diameter, surrounding the dwarf nova AT Cancri. This second discovery demonstrates that nova shells surrounding Z Cam-type dwarf novae cannot be very rare. The shell geometry is suggestive of bipolar, conical ejection seen nearly pole-on. A spectrum of the brightest AT Cnc shell knot is similar to that of the ejecta of the classical nova GK Per, and of Z Cam, dominated by [NII] emission. Galex FUV imagery reveals a similar-sized, FUV-emitting shell. We determine a distance of 460 pc to AT Cnc, and an upper limit to its ejecta mass of $\\sim  5$ $\\times 10^{-5} M_\\sun$, typical of classical novae. ",
        "introduction": "Dwarf and classical novae are all close binary stars, wherein a white dwarf accretes hydrogen-rich matter from its Roche-lobe filling companion, or from the wind of a nearby giant. In dwarf novae, an instability \\citep{osa74} episodically dumps much of the accretion disk onto the white dwarf. The liberation of gravitational potential energy then brightens these systems by up to 100-fold every few weeks or months \\citep{war95}. This accretion process in dwarf novae must inevitably build an electron degenerate, hydrogen-rich envelope on the white dwarf \\citep{sha86}. Theory and detailed simulations predict that once the accreted mass M$_{env}$ reaches of order $10^{-5}M_\\odot$, a thermonuclear runaway (TNR) will occur in the degenerate layer of accreted hydrogen. The TNR causes the rapid rise to $\\sim 10^{5} L_\\sun$ or more, and the high-speed ejection of the accreted envelope (\\citet{sha89} and \\citet{yar05}) in a classical nova explosion. Theory thus predicts that dwarf novae must inevitably give rise to classical novae.  \\citet{col09} have updated the seminal work of \\citet{rob75}, finding no evidence for dwarf nova eruptions in the progenitors of classical novae during the decades before the nova explosions. The identified progenitors of almost all classical novae are, instead, novalike variables, in which the mass transfer rate $\\dot{M}$ through the accretion disk is too high to permit the disk instability that drives dwarf nova eruptions. This apparent contradiction with theory is explained by the hibernation scenario of cataclysmic variables \\citep{sha86} as follows.  During the millenia before a nova eruption, gravitational radiation drives the white and red dwarfs closer together, enhancing Roche lobe overflow and $ \\dot{M}$. The increasingly high mass transfer rate turns a dwarf nova (DN) into a novalike variable centuries before the envelope mass reaches the value needed for a TNR. The higher $\\dot{M}$ of a novalike variable chokes off dwarf nova eruptions, hence none are seen as nova progenitors. During the few centuries after a nova eruption the mass transfer rate remains high (due to irradiation of the red dwarf), which again prevents dwarf nova eruptions. A few centuries after a nova eruption, the hibernation scenario predicts that dwarf nova eruptions should begin anew. This is because irradiation of the red dwarf by the cooling white dwarf drops, as does $\\dot{M}$. These newly reborn DN  will be the highest mass transfer rate dwarf novae - the Z Camelopardalis stars. Thus within the context of the hibernation scenario one expects old novae to evolve from novalike variables into Z Cam stars in the centuries after nova eruptions. Only these Z Cam stars will be surrounded by old nova shells. As gravitational radiation eventually drives the two stars in a CV together, one expects Z Cam stars to be the most likely progenitors of the novalike variables before they erupt as classical novae. These Z Cam stars will not be surrounded by old nova shells - their shells dispersed many millenia ago. The hibernation scenario thus predicts that some, but not all Z Cam stars should be surrounded by old nova shells.  In 2007 we reported the discovery of a classical nova shell surrounding the prototypical dwarf nova Z Camelopardalis \\citep{sha07}. This shell is an order of magnitude more extended than those detected around any other classical nova.The derived shell mass matches that of classical novae, and is inconsistent with the mass expected from a dwarf nova wind or a planetary nebula. The Z Cam shell observationally linked, for the first time, a prototypical dwarf nova with an ancient nova eruption and the classical nova process. This was the first-ever confirmation of a key prediction of cataclysmic binary TNR theory: the accreting white dwarfs in dwarf novae must eventually erupt as classical novae.  Motivated by this discovery, we have been searching for other nova shells surrounding dwarf novae. One of our targets was the Z Cam-like dwarf nova AT Cancri. In a study of AT Cnc, \\citet{bon74} found shallow, broad absorption lines, and suggested that the star is an eclipsing binary composed of a DA white dwarf and a faint red dwarf companion. \\citet{nog99} found the orbital period to be 0.2011 days,  and detected P Cygni profiles in the asymmetric $H\\alpha$ line. A summary of AT Cnc's properties, as well as spectroscopic and photometric observations are given by \\citet{nog99}.  Our optical narrowband imaging of AT Cnc immediately revealed fragmented rings surrounding the star. Followup observations with GALEX confirmed the presence of FUV - emitting material surrounding this dwarf nova.  In Section 2 we describe our observations. We show optical narrowband imagery of the rings of material surrounding AT Cancri in Section 3, and a spectrum of the shell material in section 4. GALEX ultraviolet imagery is presented in section 5. We determine the distance to AT Cnc, and an upper limit to its ejecta mass in section 6. The age of the AT Cnc ejecta is discussed in section 7. The implications of the existence of the ejecta are considered in Section 8, and we briefly summarize our results in Section 9.  \\section {Observations and Image Processing}  Narrowband images of AT Cnc in the lines of H$\\alpha$ and [NII], and broadband R images  were obtained with the 90Prime camera \\citep{wil04} of the 2.3 meter Steward Observatory telescope on 11 November 2007. The camera's focal plane array is populated with a mosaic of four thinned Lockheed 4096 x 4096 pixel CCDs. The camera provides a plate scale of 0.45 arcsec per pixel and a total field-of-view of 1.16 degree x 1.16 degree.The R images' total exposure time was 1800 seconds, while the H$\\alpha$ +  [NII] images totaled 5400 seconds. Followup imagery in the same filters was obtained with the Mosaic CCD camera at the prime focus of the Kitt Peak National Observatory Mayall 4 meter telescope. The Mosaic camera on the 4 meter telescope has eight 2048 x 4096 SITe thinned CCDs, and an image scale of 0.26 arcsec/pixel.  Imaging was carried out on the nights of 07 and 09 February 2010, and conditions were generally clear. 18 R band images, each of 180 seconds duration (3240 seconds total exposure) and 18 H$\\alpha$ +[NII] images of 1800 seconds each (32400 seconds exposure) were obtained. Images were dithered over both nights during each epoch.  After flatfielding and de-biassing, standalone Daophot \\citep{stn87} was used to align the images on each chip; then all of the chips were matched together.  All of the continuum (hereafter \"R\") and all of the narrowband (hereafter \"[NII]\") images of each epoch were combined to create the deepest possible image.  The images were stitched together using montage2, a mosaicking program within the standalone Daophot.  After this process was completed individually for both the [NII] and R band images, the narrow and broadband images were matched up with Daophot (which uses triangular stellar patterns for its matching algorithm).  Spectra of the brightest knots in the AT Cnc ejecta were obtained with the R-C spectrograph of the Kitt Peak National Observatory Mayall 4 meter telescope.  A 158 l/mm grating was used because of the faintness of the nebulosity. We combined, using IRAF imcombine, four images (1 x 600 sec plus 3 x 1800 sec, for total of 6000 sec) for the first slit position and two images (1800 sec x 2, for a total of 3600 sec) for the second slit position.  We extracted spectra of each knot using the IRAF task apall.  Ultraviolet imagery was also obtained with the NASA GALEX satellite. The GALEX image data include far-UV (FUV; $\\lambda_{eff}$=1516~\\AA, $\\Delta\\lambda$=256~\\AA) and near-UV (NUV; $\\lambda_{eff}$=2267~\\AA, $\\Delta\\lambda$=730~\\AA) images in circular fields of diameter $1\\fdg2$. The total exposure in the FUV filter is 1600 sec while that in the NUV filter is 12100 sec. The spatial resolution is $\\sim$5\". Details of the GALEX instrument and data characteristics can be found in \\citet{mar05} and \\citet{mor05}. The imaging data have been processed under the standard GALEX survey pipeline.  \\section {Imaging of the AT Cnc Shell}  The resulting R and net H$\\alpha$+[NII] (narrowband minus R) images from the KPNO 4 meter telescope, taken in 2010 are shown in Figures 1 and 2, respectively. At Cnc is circled in both images. This is one of the deepest narrowband-broadband image pairs ever taken of any nova ejecta. The net narrowband image is dominated by the striking arcs running from the NE through N, W and S of the central star. The two arcs, and their geometry are reminiscent of the hourglass-shaped nebulae of the LMC supernova SN87A \\citep{law00} and the planetary nebula  MyCn 18 \\citep{sah99}. The near coincidence of the two rings suggests that we are viewing the hourglass almost along its long symmetry axis. Both the rings are extremely fragmented, like the central ring of SN87A seen a decade and more after the supernova eruption. This morphology is suggestive of a fast wind colliding with slow or stationary ejecta from previous outbursts. The shock-dominated spectrum of the brightest knots (next section) supports this interpretation.  \\section {Spectrum of the AT Cnc Ejecta}  We placed the 4-meter spectrograph slit at two orientations and positions in order to get spectra of the two brightest knots in AT Cnc's shell.  The first slit, at an angle of 6.6 degrees from North, spans two bright knots very close to each other.  The second slit, at an angle of 119.9 degrees from North, spans a single bright knot.  The slit positions can be seen overlaid on the image of AT Cnc and its shell in Figure 2. Only the first (6000 sec spectrum) yielded sufficient signal to clearly identify the primary emission lines. That spectrum is shown in Figure 3, which is dominated by the emission lines of [NII], [OIII] and [OII].  The presence of the [OII] places an upper limit on the density of the emitting gas of about 3000 $cm^{-3}$ \\citep{app88}. Unfortunately our spectral resolution is too low to resolve the [OII] doublet and further constrain the density. As we show in section 6, AT Cnc's luminosity (roughly 1 L$_\\odot$) and effective temperature (about 10 kKelvin) are too low to photoionize the arcs of ejecta seen in Figure 2. The lack of Balmer lines and the presence of strong [NII] lines (see figure 4) suggests a shock temperature in excess of 20 kKelvins. The emission lines and their ratios are reminiscent of the spectra of the ejecta of the classical nova GK Per (nova 1901)\\citep{szd12} and the recurrent nova T Pyx \\citep{con97}. Both these are shock ionized due to the collision of rapidly outflowing ejecta with slower moving matter. ",
        "conclusions": "We report the optical narrowband detection of fragmented rings, 3 arcmin in diameter, surrounding the dwarf nova AT Cancri. The shell geometry is suggestive of bipolar, conical ejection seen nearly pole-on. A spectrum of the brightest part of the AT Cnc ejecta is similar to that of the ejecta of the classical nova GK Per, and of Z Cam, dominated by [NII], [OII] and [OIII] emission. The ejecta must be shock ionized. Galex FUV imagery reveals a similar-sized, FUV-emitting shell. We determine that AT Cnc is about 460 pc from Earth, with a system luminosity at maximum brightness that is $\\sim  L_\\sun$. The 1.5 arcmin radius of the AT Cnc ejecta corresponds to 0.2 pc at that distance, with a maximum shell mass of  $\\sim  5 \\times 10^{-5} M_\\sun$, in excellent agreement with theoretical predictions of nova ejecta masses."
    },
    "1208/1208.4693_arXiv.txt": {
        "abstract": "The structure and dynamics of the solar corona is dominated by the magnetic field. In most areas in the corona magnetic forces are so dominant that all non-magnetic forces like plasma pressure gradient and gravity can be neglected in the lowest order. This model assumption is called the force-free field assumption, as the Lorentz force vanishes. This can be obtained by either vanishing electric currents (leading to potential fields) or the currents are co-aligned with the magnetic field lines. First we discuss a mathematically simpler approach that the magnetic field and currents are proportional with one global constant, the so-called linear force-free field approximation. In the generic case, however, the relation between magnetic fields and electric currents is nonlinear and analytic solutions have been only found for special cases, like 1D or 2D configurations. For constructing realistic nonlinear force-free coronal magnetic field models in 3D, sophisticated numerical computations are required and boundary conditions must be obtained from measurements of the magnetic field vector in the solar photosphere. This approach is currently of large interests, as accurate measurements of the photospheric field become available from ground-based (for example SOLIS) and space-born (for example Hinode and SDO) instruments.  If we can obtain accurate force-free coronal magnetic field models we can calculate the free magnetic energy in the corona, a quantity which is important for the prediction of flares and coronal mass ejections. Knowledge of the 3D structure of magnetic field lines also help us to interpret other coronal observations, e.g., EUV-images of the radiating coronal plasma. ",
        "introduction": "\\label{section:introduction}  The magnetic activity of the Sun has a high impact on Earth. As illustrated in Figure~\\ref{soho1}, large coronal eruptions like flares and coronal mass ejections can influence the Earth's magnetosphere where they trigger magnetic storms and cause aurorae. These coronal eruptions have also harmful effects like disturbances in communication systems, damages on satellites, power cutoffs, and unshielded astronauts are in danger of life-threatening radiation.% \\epubtkFootnote{For an animation of a coronal mass ejection (CME) causing a substorm and aurora, see http://sohowww.nascom.nasa.gov/gallery/Movies/recon/reconsm.mpg } The origin of these eruptive phenomena in the solar corona is related to the coronal magnetic field as magnetic forces dominate over other forces (like pressure gradient and gravity) in the corona. The magnetic field, created by the solar dynamo, couples the solar interior with the Sun's surface and atmosphere. Reliable high accuracy magnetic field measurements are only available in the photosphere. These measurements, called vector magnetograms, provide the magnetic field vector in the photosphere.  \\epubtkImage{soho}{% \\begin{figure}[htb] \\centerline{\\includegraphics[width=\\textwidth]{soho1}} \\caption{Magnetic forces play a key role in solar storms that can impact Earth's magnetic shield (magnetosphere) and create colorful aurora. [Source: http://sohowww.nascom.nasa.gov/gallery/images/magnetic\\_clean.html. Courtesy of SOHO consortium. SOHO is a project of international cooperation between ESA and NASA.]} \\label{soho1} \\end{figure}}  To get insights regarding the structure of the coronal magnetic field we have to compute 3D magnetic field models, which use the measured photospheric magnetic field as the boundary condition. This procedure is often called ``extrapolation of the coronal magnetic field from the photosphere.'' In the solar corona the thermal conductivity is much higher parallel than perpendicular to the field so that field lines may become visible by the emission at appropriate temperatures. This makes in some sense magnetic field lines visible and allows us to test coronal magnetic field models. In such tests 2D projection of the computed 3D magnetic field lines are compared with plasma loops seen in coronal images. This mainly qualitative comparison cannot guarantee that the computed coronal magnetic field model and derived quantities, like the magnetic energy, are accurate. Coronal magnetic field lines which are in reasonable agreement with coronal images are, however, more likely to reproduce the true nature of the coronal magnetic field.  \\epubtkImage{beta}{% \\begin{figure}[htbp] \\centerline{\\includegraphics[scale=1.0]{beta}} \\caption{Plasma $\\beta$ model over active regions. The shaded area corresponds to magnetic fields originating from a sunspot region with 2500~G and a plage region with 150~G. The left and right boundaries of the shaded area are related to umbra and plage magnetic field models, respectively. Atmospheric regions magnetically connected to high magnetic field strength areas in the photosphere naturally have a lower plasma $\\beta$. The original figure was published as Figure~3 in \\cite{gary01}.} \\label{figbeta} \\end{figure}}  To model the coronal magnetic field $\\mathbf{B}$ we have to introduce some assumptions. It is therefore necessary to get some a priori insights regarding the physics of the solar corona. An important quantity is the plasma $\\beta$ value, a dimensionless number which is defined as the ratio between the plasma pressure $p$ and the magnetic pressure, \\begin{equation} \\beta=2 \\mu_0 \\frac{p}{B^2}. \\end{equation} Figure~\\ref{figbeta} from \\cite{gary01} shows how the plasma $\\beta$ value changes with height in the solar atmosphere. As one can see a region with $\\beta \\ll 1$ is sandwiched between the photosphere and the upper corona, where $\\beta$ is about unity or larger. In regions with $\\beta \\ll 1$ the magnetic pressure dominates over the plasma pressure (and as well over other non-magnetic forces like gravity and the kinematic plasma flow pressure). Here we can neglect in the lowest order all non-magnetic forces and assume that the Lorentz force vanishes. This approach is called the {\\textit{force-free field approximation} and for static configurations it is defined as: \\begin{eqnarray} \\mathbf{j}\\times\\mathbf{B} & = & \\mathbf{0}, \\label{forcebal}\\\\ \\mathbf{j} & = & \\frac{1}{\\mu_0}\\nabla \\times \\mathbf{B} \\;\\; { \\rm is \\, the \\, electric \\, current \\, density} \\label{ampere}, \\\\ \\nabla\\cdot\\mathbf{B} & = & 0 \\label{solenoidal}, \\end{eqnarray} or by inserting Equation~(\\ref{ampere}) into (\\ref{forcebal}): \\begin{eqnarray} (\\nabla \\times \\mathbf{B}) \\times\\mathbf{B} & = & \\mathbf{0}, \\label{eq:force-free} \\\\ \\nabla\\cdot\\mathbf{B} & = & 0 .\\label{eq:solenoidal} \\end{eqnarray} Equation~(\\ref{eq:force-free}) can be fulfilled either by: \\begin{equation} \\nabla \\times \\mathbf{B}=0 \\; \\quad \\mbox{current-free or potential magnetic fields} \\end{equation} or  by \\begin{equation} \\mathbf{B} \\parallel \\nabla \\times \\mathbf{B} \\; \\quad \\mbox{force-free fields}. \\end{equation}  Current free (potential) fields are the simplest assumption for the coronal magnetic field. The line-of-sight (LOS) photospheric magnetic field which is routinely measured with magnetographs are used as boundary conditions to solve the Laplace equation for the scalar potential $\\phi$, \\begin{equation} \\Delta \\phi =0 \\end{equation} where the Laplacian operator $\\Delta $ is the divergence of the gradient of the scalar field and \\begin{equation} \\mathbf{B} = -\\nabla \\phi . \\end{equation} When one deals with magnetic fields of a global scale, one usually assumes the so-called ``source surface'' (at about $2.5$ solar radii where all field lines become radial): See, e.g., \\citet{schatten69} for details on the potential-field source-surface (PFSS) model. Figure~\\ref{potentialglobal} shows such a potential-field source-surface model for May 2001 from \\cite{wiegelmann:etal04:conf}.  Potential fields are popular due to their mathematical simplicity and provide a first coarse view of the magnetic structure in the solar corona. They cannot, however, be used to model the magnetic field in active regions precisely, because they do not contain free magnetic energy to drive eruptions. Further, the transverse photospheric magnetic field computed from the potential-field assumption usually does not agree with measurements and the resulting potential field lines do deviate from coronal loop observations. For example, a comparison of global potential fields with TRACE images by \\cite{schrijver:etal05a} and with stereoscopically-reconstructed loops by \\cite{sandman:etal09} showed large deviations between potential magnetic field lines and coronal loops.  \\epubtkImage{potentialglobal}{% \\begin{figure}[htbp] \\centerline{\\includegraphics[scale=0.4]{potentialglobal}} \\caption{Global potential field reconstruction. The original figure was published in \\cite{wiegelmann:etal04:conf}.} \\label{potentialglobal} \\end{figure}}  The $\\mathbf{B} \\parallel \\nabla \\times \\mathbf{B}$ condition can be rewritten as \\begin{eqnarray} \\nabla \\times \\mathbf{B} & = & \\alpha \\mathbf{B} \\label{amperealpha}, \\\\ \\mathbf{B} \\cdot \\nabla \\alpha &=& 0 \\label{Bgradalpha}, \\end{eqnarray} where $\\alpha$ is called the force-free parameter or force-free function. From the horizontal photospheric magnetic field components $(B_{x0}, \\, B_{y0})$ we can compute the vertical electric current density \\begin{equation} \\mu_0 j_{z0}=\\frac{\\partial B_{y0}}{\\partial x}-\\frac{\\partial B_{x0}}{\\partial y} \\label{j_photo} \\end{equation} and the corresponding distribution of the force-free function $\\alpha(x,y)$ in the photosphere \\begin{equation} \\alpha(x,y)=\\mu_0 \\frac{j_{z0}}{B_{z0}} . \\label{alpha0direct} \\end{equation} Condition~(\\ref{Bgradalpha}) has been derived by taking the divergence of Equation~(\\ref{amperealpha}) and using the solenoidal condition~(\\ref{solenoidal}). Mathematically Equations~(\\ref{amperealpha}) and (\\ref{Bgradalpha}) are equivalent to Equations~(\\ref{forcebal})\\,--\\,(\\ref{solenoidal}). Parameter $\\alpha$ can be a function of position, but Equation~(\\ref{Bgradalpha}) requires that $\\alpha$ be constant along a field line. If $\\alpha$ is constant everywhere in the volume under consideration, the field is called linear force-free field (LFFF), otherwise it is nonlinear force-free field (NLFFF). Equations~(\\ref{amperealpha}) and (\\ref{Bgradalpha}) constitute partial differential equations of mixed elliptic and hyperbolic type. They can be solved  as a well-posed boundary value problem by prescribing the vertical magnetic field and for one polarity the distribution of $\\alpha$ at the boundaries. As shown by \\cite{bineau72} these boundary conditions ensure the existence and unique NLFFF solutions at least for small values of $\\alpha$ and weak nonlinearities. \\cite{boulmezaoud:etal00} proved the existence of solutions for a simply and multiply connected domain. As pointed out by \\cite{aly07} these boundary conditions disregard part of the observed photospheric vector field: In one polarity only the curl of the horizontal field (Equation (\\ref{j_photo})) is used as the boundary condition, and the horizontal field of the other polarity is not used at all. For a general introduction to complex boundary value problems with elliptic and hyperbolic equations we refer to \\cite{kaiser00}.  Please note that high plasma $\\beta$ configurations are not necessarily a contradiction to the force-free condition \\citep[see][for details]{neukirch05}. If the plasma pressure is constant or the pressure gradient is compensated by the gravity force $(\\nabla p = -\\rho \\nabla \\Psi$,  where $\\rho$ is the mass density and $\\Psi$ the gravity potential of the Sun.) a high-$\\beta$ configuration can still be consistent with a vanishing Lorentz force of the magnetic field. In this sense a low plasma $\\beta$ value is a sufficient, but not a necessary, criterion for the force-free assumption. In the generic case, however, high plasma $\\beta$ configurations will not be force-free and the approach of the force-free field is limited to the upper chromosphere and the corona (up to about $2.5\\,R_{\\rm s}$).  \\newpage ",
        "conclusions": "In this review we tried to give an overview of force-free magnetic fields, particularly model assumptions, which are important for understanding the physics of the solar corona. While the underlying mathematical equations describe stationary states and look relatively simple, solving them is by no means a trivial problem because of the nonlinear nature of the problem. Exact solutions are only available for further simplifications, like linearizing the equations or to restrict to 1D/2D for the nonlinear case. For force-free configurations in 3D, we know that (for given flux distributions in the photosphere) the magnetic field energy is bounded from below by a potential field. An upper-limit for the energy is more difficult to obtain. While the Aly--Sturrock conjecture (Section~\\ref{aly-sturrock}) claims that the upper limit is for the configurations with all magnetic field lines open, \\cite{choe:etal02} constructed solutions with energies above the Aly--Sturrock limit. These configurations contain discontinuities and the debate of the validity of the Aly--Sturrock limit is ongoing \\citep{hu04,wolfson:etal12}.  For practical computations of the 3D-field in the solar corona, one has to use numerical computations and several codes have been developed, compared, and applied. As input these codes require measurements of the magnetic field vector in the solar photosphere. However, the transverse field component contains an ambiguity in the azimuth, which has to be resolved before the data can be used for coronal magnetic field modeling. The accuracy of photospheric measurements is lower for the transverse field component compared with the line-of-sight field, and in weak field regions measurements and azimuth ambiguity removal are less trustworthy. Consequently the majority of coronal force-free field models are carried out in active regions, although methods for full-disk computations have been developed too. A further complication of using photospheric measurements as the boundary condition for force-free extrapolations is that the photospheric magnetic field is not necessarily consistent with the force-free assumption. Possible solutions are to use only the vertical magnetic field and the vertical electric current as boundary conditions, as done for the Grad--Rubin approach, to preprocess the photospheric measurements with the aim to make them compatible with force-free and other physical requirements, or to allow changes of the transverse magnetic field during the iteration of a force-free field. The latter approach has been implemented in the optimization approach and allows us to take measurement errors into account.  A major source for future research on force-free fields is SDO/HMI, which measures the photospheric magnetic field vector on the full disk, which in principle allows us to compute global coronal models as well as selecting appropriate isolated active regions with a sufficiently large field-of-view. Research on Stokes inversion, azimuth ambiguity removal, and force-free modeling for SDO/HMI data is ongoing. Another important aspect on coronal modeling is the comparison of force-free models as extrapolated from photospheric measurements with coronal images as observed, for example, with the Atmospheric Imaging Assembly \\citep[AIA;][]{lemen:etal12} on SDO. On the one hand, such a comparison is important to validate the models \\citep[see][for details]{derosa:etal09}, and, on the other hand, the 3D models help to interpret the observations. With the 3D structure of magnetic loops from the models in hand, one has important tools for modeling of plasma loops, and may gain understanding of coronal heating and plasma flows along the loops. Further steps on the research of eruptive phenomena like flares and CMEs are planned with time-dependent MHD simulations. Force-free models are planned to be used as initial equilibria, which are disturbed by photospheric plasma flows (which can be deduced, e.g., from measurements with SDO/HMI). The temporal evolution and the potential occurrence of eruptions can be investigated with ideal or resistive MHD simulations in comparison with observations. Questions are if or to which extent the configurations remain approximately force-free during eruptions, the role of thin current sheets and discontinuities, and the energy and helicity content. We aim to report about the progress in these aspects in an update of this review in due time."
    },
    "1208/1208.4370_arXiv.txt": {
        "abstract": "In May 2012 two asteroids made near-miss ``grazing\" passes at distances of a few Earth-radii: 2012 KP24 passed at nine Earth-radii and 2012 KT42 at only three Earth-radii. The latter passed inside the orbital distance of geosynchronous satellites. From spectral and imaging measurements using NASA's 3-m Infrared Telescope Facility (IRTF), we deduce taxonomic, rotational, and physical properties. Their spectral characteristics are somewhat atypical among near-Earth asteroids: C-complex for 2012 KP24 and B-type for 2012 KT42, from which we interpret the albedos of both asteroids to be between 0.10 and 0.15 and effective diameters of $20\\pm2$ and $6\\pm1$ meters, respectively. Among B-type asteroids, the spectrum of 2012 KT42 is most similar to 3200 Phaethon and 4015 Wilson-Harrington. Not only are these among the smallest asteroids spectrally measured, we also find they are among the fastest-spinning: 2012 KP24 completes a rotation in $2.5008\\pm0.0006$ minutes and 2012 KT42 rotates in $3.634\\pm0.001$ minutes. ",
        "introduction": "\\label{sec:Introduction}  Near-Earth asteroids (NEAs) that graze the Earth at distances of a few Earth-radii are a potential source of concern, but also present an opportunity to study small-sized asteroids in our immediate neighborhood. During a short window of opportunity, astronomers can measure the spectrum and brightness variation, from which the composition, size, rotation period, and shape can be interpreted.  The results we report come from observations we conducted as part of a rapid response program on NASA's Infrared Telescope Facility (IRTF) on Mauna Kea, Hawaii.  Our program aims to measure the spectra of these Earth-grazing asteroids in near-infrared wavelengths. Due to the small sizes of these objects they are discovered only a few hours to a few days before encounter time forcing us to execute a target of opportunity program on short notice.  Here we report on two encounters, separated by a day only. The first object, 2012 KP24, crosses the orbits of the Earth and Mars and was discovered by the Catalina Sky Survey (Drake et al. 2009) five days before (geocentric) closest approach at 9 Earth-radii. The second object, 2012 KT42, was discovered by the Catalina Sky Survey only 23 hours before its closest approach at 3 Earth-radii, bringing it within the orbital distance of geosynchronous satellites (6.6 Earth-radii). The orbital elements of these two NEAs differ from one another and they do not have any physical relationship. ",
        "conclusions": "\\label{sec:discussion}  The effective diameters of these objects are estimated from their absolute magnitudes {\\it H} and inferred albedo values ${\\it P_V}$. We use the absolute magnitude that appears in the Minor Planet Center (MPC): 26.4 mag for 2012 KP24 and 29 mag for 2012 KT42. These values, represent the magnitude of an asteroid in zero phase angle, are estimated by the MPC based on measurements taken on specific phase angles and a {\\it G} slope of 0.15. To estimate the uncertainty of the absolute magnitude, {\\it dH}, we calculate possible extreme values of the absolute magnitude by using extreme values of {\\it G} slope 0 and 0.5 (Reddy et al. 2012), and by using the maximal phase angle the asteroid was observed at (70 degrees for 2012 KP24 and 6 degrees for 2012 KT42). This derives {\\it dH} of +0.6 / -0.4 for 2012 KP24 and {\\it dH} of +0.2 / -0.1 for 2012 KT42 (the error for 2012 KP24 is larger because it was observed at higher phase angle).  Based on the derived classes there is high likelihood the objects have relatively low albedos. However, the absence of any thermal flux (causing an upturn in the 2.0-2.5 micron region on the spectra of NEAs; Rivkin et al. 2005, Reddy et al. 2012) seems counter-indicative. This contradiction can be solved when assuming a regolith-free surface (typical of small-sized asteroids; Delbo et al. 2007) that increase the \u201cbeaming parameter\u201d, $\\eta$, of the thermal model, and reduce the thermal excess to insignificant values (see Fig. 4 in Rivkin et al. 2005). Therefore, we can infer only that the asteroid albedos are no lower than 0.10 and most likely for these types in the range 0.10 to 0.15. This range is consistent with the average albedo values of these taxonomies measured by the Wide-field Infrared Survey Explorer (WISE; Mainzer et al. 2011) and the value for 3200 Phaethon (Dumas et al. 1998). We estimate the maximum effective diameters, $D_{eff}$, using: \\begin{equation} D_{eff} = \\frac{1329}{\\sqrt{P_V}}\\times10^{-H/5}. \\label{eq:Deff} \\end{equation} The uncertainty of the diameter is calculated using the extreme values of the absolute magnitude and the albedo. We derive maximum diameter values of $20\\pm6$ meters for 2012 KP24 and $6\\pm1$ meters for 2012 KT42. We deliberately avoid stating a minimal limit for the diameter, because there is a possibility that the albedos of the two asteroids are higher. However, such a minimal limit is still within a factor of two from the given value of the effective diameter.  The sizes and lightcurve results make the two studied objects among the smallest and fastest-spinning measured asteroids. We mark the asteroids on the diagram of diameter vs. spin rate (Figures~\\ref{fig:DiamSpin}) to display how the observation of such Earth grazing asteroids can help map the almost uncharted area of small, fast, and presumably intact asteroidal fragments that are typical of meteoroids capable of delivering samples. Excluding 2008 TC3, that crashed on Earth on October 2008 and its low signal-to-noise spectrum in the visible regime appears relatively flat (Jenniskens et al. 2009), none of the other seven objects among these asteroids ($D_{eff} \\leq 200$ m, fast rotation and a known taxonomy) are of C-type or B-type as 2012 KP24 and 2012 KT42; implying that monolithic-structured asteroids can be formed from dark-type composition as well and they are not necessarily composed out of S-, X- or V-type materials.  Moreover, systematic measurement of these small and fast-rotating asteroids will determine if there is an observational bias against slow-rotators of the same size. A non-bias scenario will strengthen the idea that the Yarkovsky-O'Keefe-Radzievskii-Paddack (YORP) effect, which applies a thermal torque due to the re-emission of sunlight from an asymmetric surface (Rubincam 2000), has a fast and effective influence on such Near-Earth asteroids. Using MORIS for photometry while it is guiding the telescope to track an asteroid is an applicable method to answer this scientific question.    \\begin{figure*} \\centerline{\\includegraphics[width=17cm]{Diam-Spin_KP24_KT42_BW.eps}} \\caption{The diameter-spin rate diagram of asteroids marked with 2012 KP24 (square) and 2012 KT42 (star). The errors for these two asteroids are smaller than the symbols. The background population of asteroids is divided between Near-Earth asteroids (NEAs, grey dots) and main-belt asteroids (MBAs, black dots). Only secure rotation periods are plotted (quality code $U\\geq2$). The data for the background population was taken from the Light Curve Data Base (LCDB), version of June 2012 (Warner et al. 2009). \\label{fig:DiamSpin}} \\end{figure*}"
    },
    "1208/1208.4146_arXiv.txt": {
        "abstract": "{ We present a method for identifying localized secondary populations in stellar velocity data using Bayesian statistical techniques. We apply this method to the dwarf spheroidal galaxy Ursa Minor and find two  secondary objects in this satellite of the Milky Way. One object is kinematically cold with a velocity dispersion of $4.25 \\pm 0.75\\ \\kms$ and centered at $(9.1\\arcmin \\pm 1.5, 7.2\\arcmin \\pm 1.2)$ in relative RA and DEC with respect to the center of Ursa Minor. The second object has a large velocity offset of $-12.8^{+1.75}_{-1.5}\\  \\kms$ compared to Ursa Minor and centered at $(-14.0\\arcmin^{+2.4}_{-5.8}, -2.5\\arcmin^{+0.4}_{-1.0})$. The kinematically cold object has been found before using a smaller data set but the prediction that this cold object has a velocity dispersion larger than $2.0\\ \\kms$ at 95\\% C.L. differs from previous work. We use two and three component models along with the information criteria and Bayesian evidence model selection methods to argue that Ursa Minor has one or two localized secondary populations. The significant probability for a large velocity dispersion in each secondary object raises the intriguing possibility that each has its own dark matter halo, that is, it is a satellite of a satellite of the Milky Way.\\\\ } {{\\sc keywords: Dark Matter: Substructure, Dwarf Galaxies: Ursa Minor, Bayesian Statistics}} ",
        "introduction": "The Milky Way dwarf spheroidal galaxies (dSphs) are the faintest but most numerous of the Galactic satellites. About 22 dSphs have been discovered with nine known before the Sloan Digital Sky Survey (SDSS). The latter satellites are often collectively referred to as the classical dSphs. Thus, thanks to the advent of the SDSS, the number of known Milky Way dSphs has more than doubled \\citep{Willman05, Belokurov06, Zucker06a, Zucker06b, Belokurov07, Sakamoto06, Irwin07, Walsh07}. The classical systems are in general brighter and more extended than their post-SDSS counterparts, usually referred to as the  ultra-faint dwarfs. The dSph population of the Milky Way have a wide range of luminosities, $10^{3-7} L_{\\odot}$, and sizes (half-light radii) from 40 to 1000 pc \\citep{Mateo98, Simon07,Martin08}, but span a narrow range of dynamical mass: $M(r < 300 \\rm pc) \\approx 10^7$ M$_{\\odot}$ for most of the dwarfs \\citep{Strigari08}. In the context of hierarchical structure formation scenario, these dSphs would reside in the dark matter subhalos of the Milky Way host halo and so the dynamical mass provides an estimate of the amount of dark matter in subhalos. The dynamical mass-to-light ratios span a large range of 8-4000 (in solar units); some of these systems are the most dark matter dominated systems known \\citep{Walker09c,Wolf10,Simon11,Martinez11}.  Simulations also predict that subhalos should have their own subhalos (``sub-subhalos\", e.g., \\citealt{Shaw07, Kuhlen08, Springel08, Diemand08}). While their presence in cold dark matter simulations has been verified, the mass function of these sub-subhalos hasn't been well-quantified. The subhalo mass function is seen to follow a universal profile when scaled to the virial mass of the host halo. If the sub-subhalos follow the same pattern, then we expect to see a sub-subhalo with $\\Vmax\\simeq 0.3 \\Vmax(\\mathrm{subhalo})$ \\citep{Springel08}.  We are motivated by this fact to search for stellar content that could be associated with these sub-subhalos.   Several dSphs show signs of stellar substructure or multiple distinct chemo-kinematic populations (Fornax, Sculptor, Sextans, Ursa Minor, Canes Venatici I). For instance, in Fornax, there are stellar over-densities along the minor axis, possibly remnants of past mergers \\citep{Coleman04, Coleman05} and five globular clusters \\citep{Mackey03}. In addition, Fornax's metal-rich and metal-poor stellar components seem to have different velocity dispersions \\citep{Battaglia06}. Similarly,  Sextans and Sculptor  each contain two kinematically distinct secondary populations with different metallicities \\citep{Bellazzini01, Battaglia08}. Sculptor's populations have different velocity dispersion profiles, in addition to their distinct metalicities \\citep{Battaglia08}, whereas Sextans has localized kinematically distinct population either near its center \\citep{Kleyna04, Battaglia11} or near its core radius \\citep{Walker06}. There are claims of two populations with distinct velocity and metallicity distributions in the brightest ultra-faint dwarf, Canes Venatici I (CVI) \\citep{Ibata06}, but this is not seen in two other data sets  \\citep{Simon07, Ural10}. The Bo\\\"{o}tes I ultra-faint could also have two kinematically distinct populations with different scale lengths \\citep{Koposov11}, although this wasn't apparent in earlier data sets \\citep{Munoz06-bootes, Martin07}. The largest of these Bo\\\"{o}tes I data sets contains 37 member stars and this has to be weighed against the results of \\citet{Ural10} who suggest that at least 100 stars are required to differentiate two populations.  \\begin{table} \\label{tab:umi} \\caption{Observed and derived properties of Ursa Minor.} \\begin{tabular}{lr} Parameter & Value \\\\ \\hline \\hline Distance$^{\\, 1}$ & $77 \\pm 4 $ kpc \\\\ Luminosity$^{\\, 1}$  & $3.9^{+1.7}_{-1.3} \\times 10^5 {\\rm L}_{\\odot,{\\rm V}}$\\\\ Core radius$^{\\, 1}$ & $17.9\\arcmin \\pm 2.1$\\\\ Tidal radius$^{\\, 1}$ & $77.9\\arcmin \\pm 8.9$\\\\ Half-light radius$^{\\, 1}$  & $0.445 \\pm 0.044$ kpc \\\\ Deprojected half-light radius$^{\\, 1}$ ($r_{1/2}$) & $0.588 \\pm 0.058$ kpc\\\\ Average velocity dispersion$^{\\, 2}$ & $11.61 \\pm 0.63 \\: \\kms$ \\\\ Mean velocity$^{\\, 2}$ & -247 km/s\\\\ Dynamical mass within $r_{1/2}$$^{\\, 1}$ & $5.56^{+0.79}_{-0.72} \\times 10^7 M_{\\odot}$ \\\\ Mass-to-light ratio within $r_{1/2}$$^{\\, 1}$ & 290$^{+140}_{-90}  M_{\\odot}/ {\\rm L}_\\odot$\\\\ Ellipticity$^{\\, 3}$ & $0.56 \\pm 0.05 $ \\\\ Center (J2000.0)$^{\\, 4}$ & $(15^h09^m10^{s}.2, 67^{\\circ}12'52'')$\\\\ Position angle$^{\\, 5}$ & $49.4^{\\circ}$\\\\ \\hline \\end{tabular}\\\\ Note: References are as follows 1. \\cite{Wolf10} and references therein 2. This paper  3. \\cite{Mateo98} 4. \\cite{Kleyna03} 5. \\cite{Kleyna98}\\\\ \\end{table}   Among the classical dSphs, only Draco has a lower V-band luminosity but Ursa Minor is twice as extended as Draco (in terms of its half-light radius) \\citep{Irwin95, Palma03}. Its observed and derived properties are summarized in Table~\\ref{tab:umi}. Ursa Minor is also likely the most massive satellite in terms of its dark matter halo, apart from the Magellanic clouds and the disrupting Sagittarius dSph. These properties make Ursa Minor an ideal target to search for substructure. The $\\Vmax$ at infall for the subhalo hosting Ursa Minor should be greater than 25 km/s but probably no larger than about 50 km/s \\citep{Boylan_Kolchin2012} and thus we can expect Ursa Minor to have a sub-subhalo with $\\Vmax$ in the range of $8-16 \\kms$. Despite its low mass, such a small sub-subhalo could have held on to its gas because it was protected by the deeper potential well of Ursa Minor.  Several photometric studies with different magnitude limits and overall extent observed, have reported additional localized stellar components of the stellar distribution that deviates from a smooth density profile \\citep{Olszewski85, Kleyna98, Palma03}, particularly near the center \\citep{Demers95, Eskridge01}. To the northwest of the center, a secondary peak in the spatial distribution is seen in contours and isopleths \\citep{Irwin95, Kleyna98, Bellazzini02, Palma03}. However, different studies have concluded that this secondary peak is inconclusive or of low significance \\citep{Irwin95, Kleyna98, Bellazzini02, Palma03}. Smaller scale stellar substructure is, however, seen with higher significance \\citep{Eskridge01, Bellazzini02}. Combining proper motion information with shallow photometric data in the central 20 arcmin of Ursa Minor, \\citet{Eskridge01} claim that the distribution of stars in Ursa Minor shows high significance for substructure in clumps of $\\sim 3\\arcmin0$ in size. \\citet{Bellazzini02} used the presence of a secondary peak in the distribution of the distance to the 200th neighboring star to argue that the surface density profile of Ursa Minor is not smooth. In addition, the stellar density is not symmetric along the major axis with the density falling more rapidly on the Western side \\citep{Eskridge01, Palma03} Statistically significant S-shaped morphology is also seen in contours of the red giant branch stars \\citep{Palma03}.   Spectroscopic studies of Ursa Minor \\citep{Hargreaves1994, Armandroff1995, Kleyna03, Wilkinson04, Munoz05} have shown a relatively flat velocity dispersion profile of $\\sigma \\approx 8-12 \\kms$. \\citet{Kleyna03} (K03) used a two component model to test whether the second peak in photometry was a real feature. They found a second kinematically distinct population with $\\sigma = 0.5 \\kms$ and $\\Delta \\ov = -1 \\kms$. Our results lends support to this discovery by K03 but we do not agree on the magnitude of the velocity dispersion of the substructure. We discuss this in greater detail later.  K03 argued through numerical simulations that the stellar clump they discovered could survive if the dark matter halo of Ursa Minor had a large core (about $0.85~\\rm kpc$)  but not a cusp like the prediction for inner parts of halos of $1/r$ from CDM simulations \\citep{Navarro97}. Similar numerical simulations including the Ursa Minor stellar clump have confirmed this result \\citep{Lora2012}. Similar conclusions have been reached using the observed projected spatial distribution of the five globular clusters in Fornax dSph \\citep{Mackey03}. The survival of these old globular clusters has been interpreted as evidence that the dark matter halo of Fornax may have a large core in stark contrast to the predictions of dark-matter-only CDM simulations  \\citep{Goerdt06, Sanchez2006, Cowsik09, Cole2012}. Thus, the study of the properties of the substructure in Ursa Minor has far reaching implications for the dark matter halo of this dSph and by extension the properties of the dark matter particle. Our study is complementary to the recent studies using the presence of multiple stellar populations in Fornax and Sculptor that also seem to point towards a cored dark matter density profile \\citep{Battaglia08, Walker2011, Amorisco2012}.  Current methods for finding kinematic substructure in the dSphs has relied on likelihood comparison parameter tests \\citep{Kleyna03, Ural10}, non-parametric Nadaraya-Watson estimator \\citep{Walker06}, or metalicity cuts and kinematics \\citep{Battaglia11}, but not Bayesian methods. \\cite{Hobson03} presented a Bayesian method for finding objects in noisy data. The object detection method is able to find two or more objects using only a two component model in photometric data. This method can be extended to include spectroscopic line-of-sight velocity data to search for objects using kinematics, as well as structural properties. We extend and apply this method to Ursa Minor to search for stellar substructure \\citep{Irwin95, Kleyna98} and the kinematically cold feature found by K03.  \\subsection{Data and Motivation for more Complex Models}  \\begin{figure*} \\includegraphics[width=83mm]{UmiBinGauss2.ps} \\hspace{2mm} \\includegraphics[width=83mm]{UmiBinGauss-everything.ps} \\caption{ The binned line-of-sight velocity data (red dashed) in Ursa Minor. {\\em Right:} Over-plotted is the most probable Gaussian with $\\sigma = 11.51$ and an $\\ov = -247.25$ (black solid) from the null model (single Gaussian component). {\\em Left:} The line-of-sight velocity distributions of the secondary objects and primary populations.The lines correspond to the velocity dispersions of different populations found with the Bayesian object detection method; velocity offset object (blue dot-dot-space), cold object (green dotted), primary distribution (purple dot-dash), and the total (black solid). Each component is weighted by its average number of stars found using the Bayesian object detection method. The additional kinematic components provide a better fit to the Ursa Minor data. } \\label{fig:vel} \\end{figure*}   The spectroscopic data used contains 212 Ursa Minor member stars \\citep{Munoz05}; the sample that K03 used to discover the cold feature contained 134 stars. Figure ~\\ref{fig:vel} (left) shows the radial velocities binned with the best fit single component Gaussian: this is a reasonable fit. The data are, however, fit better if we use a three component Gaussian model, cf., Figure~\\ref{fig:vel} (right). The mean and dispersion of these Gaussian distributions were derived from our Bayesian object detection that is the subject of this paper. As a prelude to our final results, we note that the centers of all three populations (the primary and two secondaries) found through the object detection method are spatially segregated.  Before we develop the Bayesian methodology, we would like to dissect the data to see if secondary populations are visible as strong local deviations in either mean velocity or velocity dispersion. To this end, we grid a $50' \\times 30'$ region around the center of Ursa Minor finely and for each grid point, we find the average velocity $\\ov$ and velocity dispersion $\\sigma$ in a $5' \\times 5'$ bin using the expectation-maximization (EM) method  (see Equations 12b and 13 of \\citet{Walker09b}). We disregard grid points where there are fewer than 7 stars in the bin. We have plotted the smoothed $\\sigma$ and $\\ov$ maps created using this method in Figure~\\ref{fig:contour}. The velocity dispersion map is the upper left panel and the average velocity map is the upper right panel. The data is rotated such that the major axis is aligned with the abscissa  ($\\theta = 49.4^{\\circ}$, see Table 1 for the photometric properties of Ursa Minor we use). There are two interesting features evident: in the $\\sigma$ map, roughly centered at $(11', -4')$, $\\sigma$ is significantly lower than the rest of the galaxy ($\\sigma < 6 \\kms$), and in the $\\ov$ map centered at $(-13', 6')$, the $\\ov$ significantly differs from Ursa Minor's overall average ($\\Delta |\\ov| > 10 \\: \\kms$). For reference, the entire data set has $\\sigma = 11.5 \\: \\kms$ and $\\ov = -247.2 \\: \\kms$ with the EM method and $\\sigma = 11.6 \\pm 0.6 \\: \\kms$ and $\\ov = -247.2 \\pm 0.8 \\: \\kms$ using a single component Gaussian model sampled with a Bayesian nested sampling technique (see next section for an explanation of the Bayesian methods we use). We have also plotted the number density (lower left panel) and the positions of the stars (lower right panel) in Figure~\\ref{fig:contour} to provide a sense for where the data is and how significant the features in the $\\ov$ and $\\sigma$ maps are. The number density map is created the same way as the $\\ov$ and $\\sigma$ maps and it shows that both features are in regions that are reasonably sampled. In the plot with the positions of the stars, we have also indicated the most probable locations for the centers and the extent of the the two features as found by our Bayesian object detection method.  We caution the reader that the plotted extents (tidal radii) of the these features have large error bars see Table~\\ref{tab:prior}).  The center of the dip in the velocity dispersion (upper left panel of Figure~\\ref{fig:contour})  is near the spectroscopic feature found by K03 and the secondary density peak seen in the photometry by several authors \\citep{Irwin95, Kleyna98, Bellazzini02, Palma03}. The average velocity feature we see does not correspond to any previous noted photometry or kinematic features. However, we note that the stellar isodensity contours of Ursa Minor are significantly asymmetric \\citep{Kleyna98, Palma03} and could hide both features.  Here we aim to show that these two localized kinematic features in Ursa Minor are statistically significant. We now turn to describing our Bayesian object detection method for finding secondary objects and model selection methods for assessing their significance.  \\begin{figure*} \\includegraphics[width=175mm]{hopethisworks.ps} \\caption{The local kinematics of Ursa Minor using the \\citet{Munoz05} data set. {\\em Upper Left:} A map of the velocity dispersion of Ursa Minor. A portion of the lower right quadrant drops below $6~\\kms$ while the rest of the galaxy is relatively uniform. {\\em Upper Right:} The average velocity of Ursa Minor found concurrently with the velocity dispersion. In the upper left quadrant the deviation $\\Delta \\ov > 10-15~\\kms$ relative to Ursa Minor while the rest of the galaxy does not differ more than $5~\\kms$. To make the contour plots, the velocity dispersion and the average velocity were found within a $5' \\times 5'$ bin ($5' \\simeq 110$ pc for a distance of 77 kpc). {\\em Lower Left:} The stellar density profile of the stars in the \\citet{Munoz05} data set. {\\em Lower Right:} The most probable locations and sizes (tidal radii) of the two objects using the Bayesian object detection method in Ursa Minor. Both of these locations correspond to the deviations seen in the average velocity and velocity dispersion maps. The coordinate system used here is such that the x-axis lines up with the major axis which has a position angle of 49.4$^{\\circ}$ \\citep{Kleyna98}. The adopted center for Ursa Minor was RA = $15^h 09^m 10^{s}.2$, DEC = $+67^{\\circ}12'52\"$ (J2000.0) (K03). For the entire sample, we obtain a mean velocity $\\ov =  -247.25~\\kms$ and velocity dispersion $\\sigma = 11.51~\\kms$. } \\label{fig:contour} \\end{figure*} ",
        "conclusions": "We have presented a method for finding multiple localized kinematically-distinct populations (stellar substructure) in line-of-sight velocity data. In the the nearby dwarf spheroidal galaxy Ursa Minor, we have found two secondary populations: ``cold\" and ``velocity offset (vo)\" objects. The estimated velocity dispersions are  $\\sigma_{cold} = 4.25 \\pm 0.75 \\: \\kms$ and $\\sigma_{vo} = 9.25 \\pm 1.25 \\: \\kms$, and the estimated mean velocities are $\\ov_{cold} = -246.25 \\pm 1.0 \\: \\kms$ and $\\ov_{vo} = -258.0 \\pm 1.5  \\: \\kms$. They are located at $(0.25_{-0.06}^{+0.04}, -0.07_{-0.07}^{+0.03})\\ \\rm kpc$ (cold object) and $(-0.24 \\pm 0.09, 0.23 \\pm 0.02)\\ \\rm kpc$ (velocity offset object) with respect to the center of Ursa Minor. The location of the cold object matches that found earlier by \\citet{Kleyna03}, but our results reveal that the velocity dispersion of this cold object could be large with a mean value close to $4\\ \\kms$. To assess the significance of our detections, we employed the Bayes Factor and information criteria $D_{KL}$ and DIC supplemented with the analysis of mock data sets with secondary populations, null hypothesis mock data sets and scrambled data sets.  The two secondary objects have $>98.5\\%$ C.L. in all the model selection tests employed.  If the velocity dispersions are as large as our Bayesian analysis seems to indicate, then these objects are likely undergoing tidal disruption or are embedded in a dark matter halo. The two possibilities are not exclusive of each other. If these objects are dark matter dominated, this would be the first detection of a satellite of a satellite galaxy.  As emphasized by \\citet{Kleyna03} the presence of localized substructure has important implications for inner density profile of the dark matter halo of Ursa Minor. The shape of the inner profile (cusp or core) has important implications for the properties of the dark matter particle with cold dark matter model predicting a cuspy inner density profile. If the stellar substructure is hosted by its own dark matter halo, then it has further implications for dark matter models since this would likely be the smallest bound dark matter structure discovered."
    },
    "1208/1208.6233_arXiv.txt": {
        "abstract": "We have explored the interplay of star formation and AGN activity in soft X-rays (0.5-2 keV) in two samples of Seyfert 2 galaxies (Sy2s). Using a combination of low resolution CCD spectra from {\\it Chandra} and {\\it XMM-Newton}, we modeled the soft emission of 34 Sy2s  using power law and thermal models. For the 11 sources with high signal-to-noise {\\it Chandra} imaging of the diffuse host galaxy emission, we estimate the luminosity due to star formation by removing the AGN, fitting the residual emission. The AGN and star formation contributions to the soft X-ray luminosity (i.e. L$_{x,AGN}$ and L$_{x,SF}$) for the remaining 24 Sy2s were estimated from the power law and thermal luminosities derived from spectral fitting. These luminosities were scaled based on a template derived from XSINGS analysis of normal star forming galaxies. To account for errors in the luminosities derived from spectral fitting and the spread in the scaling factor, we estimated L$_{x,AGN}$ and L$_{x,SF}$ from Monte Carlo simulations. These simulated luminosities agree with L$_{x,AGN}$ and L$_{x,SF}$ derived from {\\it Chandra} imaging analysis within a 3$\\sigma$ confidence level. Using the infrared [NeII]12.8$\\mu$m and [OIV]26$\\mu$m lines as a proxy of star formation and AGN activity, respectively, we independently disentangle the contributions of these two processes to the total soft X-ray emission. This decomposition generally agrees with L$_{x,SF}$ and L$_{x,AGN}$ at the 3$\\sigma$ level. In the absence of resolvable nuclear emission, our decomposition method provides a reasonable estimate of emission due to star formation in galaxies hosting type 2 AGN. ",
        "introduction": "Supermassive black holes (SMBHs) and their parent galaxies co-evolve \\citep{kr, magorrian, FM, gebhardt, tremaine, hr}. In particular, observational and theoretical work has established a link between accreting SMBHs (active galactic nuclei, or AGN) and host galaxy star formation. Common mechanisms have been proposed for triggering star formation while fueling SMBH accretion, including galaxy mergers \\citep[e.g.][]{Sanders, Hopkins} and secular processes involving gravitational instabilities induced by spiral arms or galactic sized bars \\citep[e.g.][]{KK, Cisternas, Schawinski}.  Disentangling star formation from AGN activity becomes a necessary endeavor to investigate SMBH and host galaxy co-evolution. This separation has been explored extensively and has included identifiying optical and ultraviolet (UV) spectral signatures of starburst activity in AGN \\citep{Cid}, performing principal component analysis on infrared spectra of AGN \\citep[e.g.][]{Buchanan}, and analyzing optical and infrared (IR) diagnostics that parameterize the relative contribution of AGN to star formation, such as ratios of infrared and optical emission lines that indicate ionization field hardness \\citep{Genzel, Kewley01, Kauffmann, Treyer, me2, me4} and strength of the polycyclic aromatic hydrocarbon features \\citep{Genzel, O'Dowd, me2, me4}. Here we extend the study of the interplay between AGN activity and star formation to the X-ray regime.  Quiescent galaxies emit thermal and non-thermal X-ray emission. Hot gas energized by stellar winds and supernova explosions radiates in soft X-rays (0.5 - 2 keV) whereas X-ray binares and supernova remnants dominate the non-thermal emission above 2 keV \\citep[see][for a review]{Fabbiano}. Due to the relatively short life time of high mass X-ray binaries (HMXBs, $\\tau < 10^{7}$ yr) and the short delay between starburst and generation of soft X-rays from hot gas, hard and soft X-ray emission can trace the instantaneous star formation rate (SFR). Indeed, X-ray emission from quiescent galaxies is well correlated with radio, infrared, optical and ultraviolet SFR indicators \\citep[e.g.][]{Ranalli, Rosa, PS}. Various X-ray SFRs have thus been calibrated in the literature based on the correlation of X-ray emission with far-infrared and radio luminosities \\citep{Ranalli}, X-ray emission from HMXBs and cumulative galactic X-ray point sources with far-infrared luminosities \\citep{Persic04, Persic07} and X-ray luminosity with infrared plus ultraviolet emission \\citep{PS}. However, low mass X-ray binaries, which trace stellar mass (M$_{\\star}$) rather than star formation, also contributes to the X-ray emission and becomes the primary source of X-ray radiation in galaxies with low SFR/M$_{\\star}$ \\citep{Colbert, Lehmer}.  An active nucleus complicates using X-ray emission as a SFR indicator: accretion onto the SMBH dominates emission above 2 keV. Isolating the hard X-ray emission due to star formation then becomes prohibitively difficult when the nucleus is not resolved. However, soft X-ray emission can have comparable contributions from scattered/reflected AGN emission and thermal emission from gas associated with starburst activity when the AGN is obscured or has very low luminosity. Determining the relative contribution of AGN and star formation to the soft X-ray emission can be achieved through high resolution spectroscopy which resolves narrow emission lines. Various diagnostics can be used to distinguish between collisionally and photoionized plasma, such as the width of the radiative recombination continua \\citep{LP96} and ratios of forbidden, intercombination and resonance lines in the OVII triplet \\citep{PD00}. However, with current X-ray missions (i.e. {\\it Chandra} and {\\it XMM-Newton}) high resolution spectra are only obtainable through grating observations, necessitating long exposure times on X-ray bright sources to achieve adequate signal to noise. How well can less time-intensive low resolution CCD spectroscopy achieve the goal of disentangling starburst from AGN activity in soft X-rays?  Local obscured AGN, Seyfert 2 galaxies (Sy2s), provide an ideal laboratory to answer this question. As these sources are relatively nearby and the accretion disk is hidden behind an obscuring ``torus,'' circumnuclear starbursts are visible. Hence, both AGN and host galaxy star formation can be modeled. Levenson et al. (2004, 2005) analyzed {\\it Chandra} observations of two Sy2/starburst composites (NGC 5135 and NGC 7130) and demonstrated that a thermal component due to stellar processes was necessary to adequately model the soft X-ray emission. We expand this methodology to Sy2s in general to investigate the efficacy of modeling soft X-ray emission with a thermal component to describe host galaxy star formation.  Our analysis focuses on two samples of Sy2s: an optically selected [OIII]5007\\AA\\ sample \\citep{me1} and a mid-infrared selected 12$\\mu$m sample \\citep{12m}. Using {\\it Chandra} and {\\it XMM-Newton} observations, we decomposed the soft X-ray emission into a star-forming and an AGN component. For Sy2s observed with {\\it Chandra} that have significant host galaxy emission, we are able to resolve the point source associated with the AGN and can therefore remove it and analyze emission soley from the host galaxy, which we ascribe to star formation (L$_{x,SF}$). For the Sy2s observed with {\\it XMM-Newton} and the remaining {\\it Chandra} sources, we fit the X-ray spectra with a combination of a thermal and power-law model. To account for the presence of X-ray binaries, we scale the soft thermal and power law fluxes using star-forming galaxies from the XSINGS sample as a template (Ptak et al. 2012), obtaining soft X-ray luminosity values associated with star formation (L$_{x,SF}$) and AGN activity (L$_{x,AGN}$). We compare these estimates of L$_{x,SF}$ and L$_{x,AGN}$ with infrared (IR) spectral signatures that accurately describe star formation ([NeII]12.8$\\mu$m luminosity, L$_{[NeII]}$,  \\citet{Ho, me4} and AGN activity ([OIV]26$\\mu$m luminosity, L$_{[OIV]}$, \\citet{Rigby, DS, Mel, me2} to test the accuracy of this decomposition.  \\renewcommand{\\thefootnote}{\\arabic{footnote}} ",
        "conclusions": "We have modeled the 0.5-10 keV spectra of two homogeneous samples of Seyfert 2 galaxies to disentangle the starburst and AGN emission in soft X-rays (0.5-2 keV). Eleven of these Sy2s, observed with {\\it Chandra}, had high signal-to-noise unresolved emission after removing the AGN. We derive L$_{x,SF}$ from spectrally fitting this unresolved emission and assign L$_{x,AGN}$ to the difference between the total soft X-ray emission and L$_{x,SF}$.  The remaining 24 sources were decomposed into an AGN and star formation component by modeling the soft X-ray emission with a power law and thermal model. The luminosities of these sources were converted to L$_{x,SF}$ and L$_{x,AGN}$ values based on a scaling factor derived from XSINGS analysis on normal star forming galaxies. To account for errors on the luminosities derived from spectral fits and the scaling factor, we executed Monte Carlo simulations, assuming a Gaussian distribution of random variables centered on the best-fit values of L$_{APEC}$, L$_{pow}$ and $R$, with the spread corresponding to the errors on these parameters. Using Eqs. 4 and 5, we calculate L$_{x,SF}$ and L$_{x,AGN}$ from the 1000 simulated parameters and use the mode of this distribution as our estimate of the X-ray luminosities due to star formation and AGN activity.  Our conclusions are summarized as follows:  \\begin{itemize}  \\item The soft X-ray spectra of 10 Sy2s were well fitted by two thermal model components, indicating the presence of a two temperature gas. This result is similar to what has been observed in starbust galaxies \\citep{Dahlem, Strickland}.  \\item For the subset of 11 Sy2s with high signal-to-noise {\\it Chandra} imaging of unresolved host galaxy emission, estimates of the soft X-ray emission due to star formation and AGN from both {\\it Chandra} imaging analysis and estimates from the Monte Carlo simulations agree at the 3$\\sigma$ level. In the absence of resolved nuclear emission, scaling L$_{pow}$ and L$_{AGN}$ by the factors derived from XSINGS analysis is thus a reasonable method to estimate the AGN and star formation contributions to the soft X-ray emission.  \\item The independent decompositions of the soft X-ray luminosity into a star formation and AGN component using IR data as a proxy and scaling L$_{APEC}$ and L$_{pow}$ based on XSINGS galaxies largely agree within a 3$\\sigma$ confidence interval. Deviations of the best fit regression line from this agreement are slight and appear at higher luminosities ($>10^{41}$ erg/s/cm$^2$) in the star formation decomposition and at moderate luminosities ($10^{39} \\lesssim$ L$\\lesssim 5\\times 10^{41}$ erg/s/cm$^2$) when describing AGN emission. Though more scatter in individual sources is evident in the AGN decomposition, the star formation relationship is more consistent among both methods.  \\item Comparison of our calculated X-ray SFR using the \\citet{PS} calibration with an FIR derived SFR from \\citet{Kennicutt} for the 12$\\mu$m sample shows general agreement.  \\end{itemize}  We have demonstrated that analysis of low resolution CCD X-ray spectra can effectively disentangle emission from AGN activity and star formation in 0.5-2keV X-rays. Using the decomposition we have presented, L$_{x,SF}$ can be used to cleanly estimate the SFR in Sy2s using existing calibrations \\citep[e.g.]{Ranalli,PS}, complementing previous studies in the optical and infrared and providing a more panchromatic view of the interplay between SMBH accretion and star formation.     \\clearpage \\begin{landscape} \\begin{deluxetable}{lccc}  \\tablewidth{0pt} \\tablecaption{\\label{o3_ap}[OIII] Sample: Aperture Extraction Areas\\tablenotemark{1}} \\tablehead{  \\colhead{Galaxy} & \\colhead{$z$} & \\colhead{Aperture Radius ('')}    & \\colhead{Aperture Radius (kpc)\\tablenotemark{2}} \\\\ &               & \\colhead{PN/MOS1/MOS2} & \\colhead{PN/MOS1/MOS2} }  \\startdata  Mrk 0609                & 0.034 & 65/34/35 & 44/23/24 \\\\  IC 0486                 & 0.027 & 34/33/35 & 19/18/19 \\\\  2MASX J08035923+2345201 & 0.029 & 17/20/18 & 10/12/11 \\\\  2MASX J08244333+2959238 & 0.025 & 40/40/40 & 21/21/21 \\\\  2MASX J10181928+3722419 & 0.049 & 20/22/22 & 20/22/22 \\\\  2MASX J11110693+0228477\\tablenotemark{3} & 0.035 & 18/-/-  & 13 \\\\  CGCG 242-028            & 0.026 & 19/22/13 & 10/12/7  \\\\  SBS 1133+572            & 0.050 & 20/16/17 & 20/16/17 \\\\  Mrk 1457                & 0.049 & 19/13/13 & 18/13/13 \\\\  2MASX J11570483+5249036 & 0.036 & 35/35/35 & 26/26/26 \\\\  2MASX J12183945+4706275\\tablenotemark{4} & 0.094 & -/30/31 & 53/55 \\\\  2MASX J12384342+0927362 & 0.083 & 33/30/30 & 52/47/47 \\\\  NGC 5695                & 0.014 & 31/22/17 & 9/7/5    \\\\   \\enddata  \\tablenotetext{1}{[OIII] sources were observed with {\\it XMM-Newton} only. PN, MOS1 and MOS2 refer to the three detectors onboard {\\it XMM-Newton}.} \\tablenotetext{2}{Conversion from arc seconds to kpc is based on the cosmology of H$_0$=70 km/s/Mpc, $\\Omega_M$=0.27 and $\\Omega_\\Lambda$ = 0.73.} \\tablenotetext{3}{Only PN spectrum fit since MOS1 and MOS2 spectra suffered from low signal-to-noise.} \\tablenotetext{4}{Source fell on chip gap in PN detector so only spectra from MOS1 and MOS2 were fit.} \\end{deluxetable}  \\end{landscape}   \\clearpage  \\begin{landscape} \\begin{deluxetable}{lccccc} \\small \\tablewidth{0pt} \\tablecaption{\\label{12m_ap}12$\\mu$m Sample: Aperture Extraction Areas\\tablenotemark{1}} \\tablehead{ \\colhead{Galaxy} & \\colhead{$z$} & \\colhead{Observatory} & \\colhead{ObsID} & \\colhead{Aperture Radius('')} & \\colhead{ApertureRadius (kpc)\\tablenotemark{2}} \\\\ &                       &                 & \\colhead{PN/MOS1/MOS2}      & \\colhead{PN/MOS1/MOS2}}  \\startdata  NGC 0424        & 0.012 & {\\it XMM-Newton} & 00029242301 & 34/33/33 & 8/7/7    \\\\ &       & {\\it Chandra}    & 03146       & 15       & 3        \\\\  NGC 1144        & 0.028 & {\\it XMM-Newton} & 0312190401  & 36/33/33 & 20/18/18 \\\\  NGC 1320        & 0.009 & {\\it XMM-Newton} & 0405240201  & 37/35/35 & 6/6/6    \\\\  NGC 1386        & 0.003 & {\\it XMM-Newton} & 0140950201  & 45/40/40 & 2/2/2    \\\\ &       & {\\it Chandra}    & 04076       & 19       & 1        \\\\  NGC 1667        & 0.015 & {\\it XMM-Newton} & 0200660401  & 32/22/34 & 10/7/10  \\\\  F05189-2524     & 0.043 & {\\it XMM-Newton} & 0085640101  & 41/31/33 & 35/26/28 \\\\ &       & {\\it Chandra}    & 02034       & 40       & 34       \\\\ &       & {\\it Chandra}    & 03432       & 40       & 34       \\\\  NGC 3982        & 0.004 & {\\it XMM-Newton} & 0204651201  & 34/37/44 & 3/3/4    \\\\ &       & {\\it Chandra}    & 04845       & 34       & 3        \\\\  NGC 4388        & 0.008 & {\\it XMM-Newton} & 0110930701  & 36/35/43 & 7/7/8    \\\\ &       & {\\it XMM-Newton} & 0110930301  & 37/32/33 & 7/6/6    \\\\ &       & {\\it Chandra}    & 01619       & 36       & 7        \\\\   NGC 4501        & 0.008 & {\\it XMM-Newton} & 0112550801  & 18/22/33 & 3/4/6    \\\\ &       & {\\it Chandra}    & 02922       & 20       & 4        \\\\  TOLOLO 1238-364 & 0.011 & {\\it Chandra}    & 04844       & 20        & 5       \\\\  NGC 4968        & 0.010 & {\\it XMM-Newton} & 0002940101  & 17/18/14 & 4/4/3    \\\\ &       & {\\it XMM-Newton} & 0200660201  & 21/18/18 & 5/4/4    \\\\  M-3-34-64       & 0.017 & {\\it XMM-Newton} & 0206580101  & 59/34/38 & 21/12/14 \\\\  NGC 5135        & 0.014 & {\\it Chandra}    & 02187       & 20       & 6        \\\\  NGC 5194\\tablenotemark{3} & 0.002 & {\\it XMM-Newton} & 0112840201  & 65/42/49 & 3/2/2 \\\\ &       & {\\it XMM-Newton} & 0303420101  & -/56/49  & -/2/2 \\\\ &       & {\\it XMM-Newton} & 0303420201  & -/61/46  & -/3/2 \\\\ &       & {\\it Chandra}    & 00354       & 40       & 2     \\\\ &       & {\\it Chandra}    & 01622       & 40       & 2     \\\\ &       & {\\it Chandra}    & 03932       & 40       & 2     \\\\  NGC 5347        & 0.008 & {\\it Chandra}    & 04867       & 6        & 1        \\\\  Mrk 463         & 0.050 & {\\it XMM-Newton} & 0094401201  & 49/45/45 & 49/45/45 \\\\ &       & {\\it Chandra}    & 04913       & 40       & 40       \\\\  NGC 5506\\tablenotemark{4} & 0.006  & {\\it XMM-Newton} & 0013140101  & 60/-/-  & 9/-/-   \\\\ &        & {\\it XMM-Newton} & 0013140201  & 54/-/-  & 8/-/-   \\\\ &        & {\\it XMM-Newton} & 0201830201  & 56/-/53 & 8/-/8   \\\\ &        & {\\it XMM-Newton} & 0201830301  & 52/-/64 & 8/-/9   \\\\ &        & {\\it XMM-Newton} & 0201830401  & 55/-/55 & 8/-/8   \\\\ &        & {\\it XMM-Newton} & 0201830501  & 68/-/98 & 10/-/14 \\\\ &        & {\\it XMM-Newton} & 0554170201  & 64/-/-  & 9/-/-   \\\\ &        & {\\it XMM-Newton} & 0554170101  & 66/-/-  & 10/-/-  \\\\  Arp 220         & 0.018  & {\\it XMM-Newton} & 0101640801  & 20/21/25 & 8/8/9   \\\\ &        & {\\it XMM-Newton} & 0101640901  & 28/25/23 & 11/9/9  \\\\ &        & {\\it XMM-Newton} & 0205510201  & 34/28/22 & 13/11/8 \\\\ &        & {\\it Chandra}    & 00869       & 40       & 15      \\\\  NGC 6890        & 0.008  & {\\it XMM-Newton} & 0301151001  & 17/13/15 & 3/2/2   \\\\  IC 5063         & 0.011  & {\\it Chandra}    & 07878       & 12       & 3       \\\\  NGC 7130        & 0.016  & {\\it Chandra}    & 02188       & 17       & 5       \\\\  NGC 7172        & 0.009  & {\\it XMM-Newton} & 0147920601  & 52/41/34 & 8/7/5   \\\\ &        & {\\it XMM-Newton} & 0202860101  & 61/59/60 & 10/9/10 \\\\ &        & {\\it XMM-Newton} & 0414580101  & 53/59/55 & 8/9/9   \\\\  NGC 7582        & 0.005  & {\\it XMM-Newton} & 0112310201  & 44/41/41 & 4/4/4   \\\\ &        & {\\it XMM-Newton} & 0204610101  & 60/43/47 & 6/4/4   \\\\ &        & {\\it Chandra}    & 00436       & 43       & 4       \\\\ &        & {\\it Chandra}    & 02319       & 43       & 4       \\\\  NGC 7674        & 0.029  & {\\it XMM-Newton} & 0200660101  & 19/18/16 & 11/10/9 \\\\  \\enddata  \\tablenotetext{1}{For {\\it XMM-Newton} observations, the extraction areas for the PN, MOS1 and MOS2 detectors are listed separately.} \\tablenotetext{2}{Conversion from arc seconds to kpc is based on the cosmology of H$_0$=70 km/s/Mpc, $\\Omega_M$=0.27 and $\\Omega_\\Lambda$ = 0.73.} \\tablenotetext{3}{Source fell on chip gap on PN detector for ObsIDs 0303420101 and 0303420201; these spectra were not fit.} \\tablenotetext{4}{Dashes indicate that the spectra suffered from pile-up in that particular detector and were therefore not fit. See \\citet{me3} for details.} \\end{deluxetable} \\end{landscape}    \\clearpage \\begin{landscape} \\begin{deluxetable}{lcccccrr}  \\tablewidth{0pt} \\tablecaption{\\label{o3_apec}[OIII] Sample: APEC model parameters (solar abundance)} \\tablehead{ \\colhead{Galaxy} & \\colhead{N$_{H,1}$} & \\colhead{kT$_1$} & \\colhead{kT$_2$} & \\colhead{$\\Gamma$} & \\colhead{N$_{H,2}$}   & \\colhead{$\\chi^2$ 2 APECs} &  \\colhead{$\\chi^2$ 1 APEC} \\\\ \\colhead{ }      & \\colhead{10$^{22}$ cm$^{-2}$} & \\colhead{keV} & \\colhead{keV} & & \\colhead{10$^{22}$cm$^{-2}$} & \\colhead{DOF} & \\colhead{DOF}}  \\startdata  Mrk 0609\\tablenotemark{1} & 0.04 & 0.27$^{+0.05}_{-0.04}$ & ...& 1.77$^{+0.05}_{-0.04}$ & ... &  ... & 160 (203) \\\\  IC 0486 & $<$0.06 & $<$0.16 & ... & 1.23$^{+0.08}_{-0.07}$ & 1.00$^{+0.10}_{-0.09}$ & ... & 636 (629) \\\\  2MASX J08035923+2345201\\tablenotemark{2} & 0.61$^{+0.10}_{-0.14}$ & $<$0.12 & ... & 2.84$^{+1.05}_{-1.18}$ & 46.7$^{+22.6}_{-24.1}$ & ... & 120 (95) \\\\  2MASX J08244333+2959238\\tablenotemark{1} & 0.03 & 0.68$^{+0.06}_{-0.06}$ & $<$0.12 & 1.54$^{+0.36}_{-0.34}$ & 15.5$^{+2.8}_{-2.4}$ & 229 (186) & 263 (188)\\\\   2MASX J10181928+3722419\\tablenotemark{1} & 0.01 & 0.18$^{+0.04}_{-0.05}$  & ... &  2.63$^{+0.64}_{-0.69}$ & ... & ... & 52.1 (61) \\\\  2MASX J11110693+0228477\\tablenotemark{3}  & $<$0.62 & 0.23$^{+0.15}_{-0.09}$ & ... & 2.04$^{+1.88}_{-0.88}$ & ... & ... & 37.5 (36) \\\\  CGCG 242-028\\tablenotemark{2} & 0.69$^{+0.17}_{-0.23}$ & $<$0.15 & ... & 0.31$^{+0.46}_{-0.49}$ & ... & ... & 87.0 (90) \\\\  SBS 1133+572 & $<$0.10 & 0.824$^{+0.23}_{-0.21}$ & ... & 3.08$^{+0.61}_{-0.38}$ & 57.6$^{+45.4}_{-30.2}$  & ... & 38.1 (48) \\\\  Mrk 1457 & 0.66$^{+0.12}_{-0.12}$ & 0.14$^{+0.03}_{-0.03}$  & ... & 1.29$^{+1.37}_{-1.14}$ & 27.5$^{+15.8}_{-9.3}$ & ... & 35.2 (35) \\\\  2MASX J11570483+5249036 & $<$0.1 & 0.17$^{+0.03}_{-0.05}$ & ... & 2.85$^{+0.23}_{-0.34}$ & 83.9$^{+47.1}_{-27.9}$ & ... & 123 (76) \\\\  2MASX J12183945+4706275\\tablenotemark{1} & 0.02 & $<$0.24  & ... & 1.95$^{+0.70}_{-0.86}$ & 87.2$^{+66.8}_{-34.1}$  & ... & 15 (19) \\\\  2MASX J12384342+0927362 & $<$0.07 & $<$0.23 & ... & 2.17$^{+0.31}_{-0.22}$ & 29.3$^{+3.1}_{-2.6}$ & ... & 195 (164) \\\\  NGC 5695 & $<$0.59 & 0.24$^{+0.73}_{-0.12}$ & ... & 2.55$^{+0.68}_{-0.52}$ & ... &  ... & 75.7 (62) \\\\   \\enddata  \\tablenotetext{1}{Best fit absorption same as Galactic absorption. This parameter was then frozen to the Galactic value.} \\tablenotetext{2}{Used c-stat.} \\tablenotetext{3}{PN only, MOS1 and MOS2 had low signal-to-noise spectra.} \\end{deluxetable}  \\end{landscape}   \\clearpage \\begin{landscape} \\begin{deluxetable}{lcccccrr} \\footnotesize \\tablewidth{0pt} \\tablecaption{\\label{12m_apec}APEC model parameters (solar abundance\\tablenotemark{1})} \\tablehead{ \\colhead{Galaxy} & \\colhead{N$_{H,1}$} & \\colhead{kT$_1$} & \\colhead{kT$_2$} & \\colhead{$\\Gamma$} & \\colhead{N$_{H,2}$}   & \\colhead{$\\chi^2$ 2 APECs} &  \\colhead{$\\chi^2$ 1 APEC} \\\\ \\colhead{ }      & \\colhead{10$^{22}$ cm$^{-2}$} & \\colhead{keV} & \\colhead{keV} & & \\colhead{10$^{22}$cm$^{-2}$} & \\colhead{DOF} & \\colhead{DOF}}  \\startdata  NGC 0424 & 0.04$^{+0.03}_{-0.03}$ & 0.69$^{+0.08}_{-0.05}$ & $<$0.13 & 2.21$^{+0.21}_{-0.28}$ & 32.6$^{+8.5}_{-6.0}$ & 246 (178) & 267 (180)\\\\   NGC 1144\\tablenotemark{1} & 0.06 & 0.37$^{+0.29}_{-0.06}$ & ... & 1.91$^{+0.37}_{-0.24}$ & 47.0$^{+3.5}_{-3.2}$ & .. & 175 (149) \\\\   NGC 1320 & $<$0.08 & 0.15$^{+0.03}_{-0.03}$ & 0.73$^{+0.09}_{-0.02}$ & 2.97$^{+0.27}_{-0.21}$ & 48.0$^{+35.6}_{-15.5}$ & 233 (188) & 279 (190) \\\\   NGC 1386 & 0.04$^{+0.02}_{-0.01}$ & 0.67$^{+0.03}_{-0.03}$ & 0.13$^{+0.02}_{-0.01}$ & 2.72$^{+0.12}_{-0.14}$ & 31.4$^{+22.9}_{-11.5}$ & 398 (332) & 435 (334) \\\\   NGC 1667\\tablenotemark{2} & 0.05 & 0.33$^{+0.07}_{-0.04}$ & ... & 2.18$^{+0.34}_{-0.37}$ & ... & ... & 49.8 (38) \\\\   F05189-2524\\tablenotemark{2} & 0.02 & 0.1 & ... & 2.05$^{+0.14}_{-0.14}$ & 6.77$^{+0.44}_{-0.42}$ & ... & 501 (374) \\\\   NGC 3982\\tablenotemark{2} & 0.01 & 0.29$^{+0.03}_{-0.03}$ & ... & 2.39$^{+0.18}_{-0.15}$ & 40.3$^{+25.5}_{-16.3}$ & ... & 174 (160)  \\\\   NGC 4388 (XMM)\\tablenotemark{2} & 0.03 & 0.60$^{+0.04}_{-0.03}$ & 0.15$^{+0.02}_{-0.03}$ & 1.25$^{+0.12}_{-0.12}$ & 24.3$^{+1.1}_{-1.0}$ & 510 (495) & 580 (498) \\\\ NGC 4388 (Chandra)\\tablenotemark{2} & 0.03 & 0.60$^{+0.04}_{-0.04}$ & 0.15$^{+0.04}_{-0.02}$ & 0.38$^{+0.29}_{-0.30}$ & 25.6$^{+3.1}_{-2.9}$ & 210 (166) &  269 (168) \\\\  NGC 4501\\tablenotemark{2} & 0.03 & 0.36$^{+0.05}_{-0.03}$ & ... & 1.52$^{+0.30}_{-0.29}$ & ... & ... & 94.2 (102) \\\\  TOLOLO 1238-364\\tablenotemark{3} & $<$0.19 & 0.32$^{+0.44}_{-0.06}$ & ... & 2.45$^{+0.33}_{-0.36}$ & ... & ... & 91.6 (105) \\\\   NGC 4968\\tablenotemark{2,3} & 0.08 & 0.68$^{+0.12}_{-0.12}$ & ... & 1.72$^{+0.18}_{-0.21}$ & ... & ... & 325 (268)\\\\   M-3-34-64\\tablenotemark{2} & 0.05 & 2.85$^{+0.34}_{-0.38}$ & 0.77$^{+0.02}_{-0.02}$ & 3.24$^{+0.19}_{-0.23}$ & 53.6$^{+2.9}_{-3.6}$ & 780 (492) & 848 (493) \\\\   NGC 5135\\tablenotemark{2} & 0.05 & 0.73$^{+0.03}_{-0.03}$ & ... & 2.34$^{+0.07}_{-0.08}$ & ... & .. & 166 (114) \\\\   NGC 5194 & 0.04$^{+0.01}_{-0.01}$ & 0.18$^{+0.01}_{-0.01}$ & 0.62$^{+0.01}_{-0.01}$ & 3.18$^{+0.14}_{-0.14}$ & 10.2$^{+0.80}_{-0.75}$ & 1560 (1291) & 2021 (1293) \\\\   NGC 5347\\tablenotemark{2,3} & 0.02 & $<$0.21 & ... & 1.53$^{+0.30}_{-0.29}$ & 32.6$^{+24.1}_{-19.6}$ & ... & 69.9 (82)\\\\   Mrk 463\\tablenotemark{2} & 0.02 & 0.74$^{+0.04}_{-0.02}$ & 0.20$^{+0.06}_{-0.04}$ & 1.76$^{+0.13}_{-0.16}$ & 28.6$^{+4.5}_{-3.4}$ & 365 (316) & 392 (318) \\\\   NGC 5506\\tablenotemark{4} & 0.11$^{+0.01}_{-0.01}$ & 0.74$^{+0.05}_{-0.05}$ & ... & 1.71$^{+0.02}_{-0.01}$ & 2.69$^{+0.02}_{-0.03}$ & ... & 2646 (2380) \\\\ NGC 5506\\tablenotemark{5} & '' & 0.94$^{+0.39}_{-0.24}$ & '' & '' & '' & ''& ''  \\\\ NGC 5506\\tablenotemark{6}& 0.17$^{+0.01}_{-0.01}$ & ... & ... & 1.83$^{+0.02}_{-0.02}$ & 2.80$^{+0.01}_{-0.02}$ & ... & 3982 (3137) \\\\ NGC 5506\\tablenotemark{7} & '' & '' & '' & 1.76$^{+0.01}_{-0.00}$ & '' & '' & '' \\\\ NGC 5506\\tablenotemark{8} & 0.12$^{+0.01}_{-0.01}$ & 0.83$^{+0.03}_{-0.03}$ & '' & '' & '' & '' & '' \\\\ NGC 5506\\tablenotemark{9} & '' & 0.96$^{+0.05}_{-0.05}$ & '' & '' & '' & '' & '' \\\\   Arp 220\\tablenotemark{1,2} & 0.04 & 0.79$^{+0.04}_{-0.04}$ & ... & 0.95$^{+0.26}_{-0.24}$ & ... & ... & 309(302) \\\\   NGC 6890\\tablenotemark{3} & $<$0.22 & 0.78$^{+0.24}_{-0.19}$ & ... & 3.28$^{+0.88}_{-0.74}$ & 27.4$^{+18.4}_{-11.3}$ & ... & 164 (148) \\\\   IC 5063\\tablenotemark{2,10} & 0.06 & 0.60$^{+0.10}_{-0.11}$ & ... & 1.85$^{+0.16}_{-0.21}$ & 21.0$^{+1.1}_{-1.2}$ & ... & 198 (141) \\\\  NGC 7130\\tablenotemark{3} & 0.06$^{+0.03}_{-0.02}$ & 0.63$^{+0.03}_{-0.03}$ & ... & 2.43$^{+0.25}_{-0.18}$ & 75.4$^{+55.6}_{-38.1}$ & ... & 334 (240) \\\\  NGC 7172\\tablenotemark{2} & 0.02 & 0.28$^{+0.05}_{-0.04}$ & ... & 1.56$^{+0.02}_{-0.03}$ & 7.36$^{+0.10}_{-0.10}$ & ... & 2074 (1746) \\\\ NGC 7172\\tablenotemark{2} & '' & 0.26$^{+0.02}_{-0.02}$ & ... & 1.54$^{+0.02}_{-0.03}$ & 8.13$^{+0.11}_{-0.11}$ & '' & '' \\\\   NGC 7582 (Chandra)\\tablenotemark{2} & 0.01 & 0.72$^{+0.05}_{-0.05}$ & ... & 1.94$^{+0.11}_{-0.10}$ & 24.6$^{+1.8}_{-1.6}$ & ... & 355 (301) \\\\ NGC 7582 (XMM)\\tablenotemark{2} & 0.01 & 0.18$^{+0.02}_{-0.04}$ & 0.72$^{+0.01}_{-0.01}$ & 1.75$^{+0.05}_{-0.03}$ & 27.0$^{+1.5}_{-1.5}$ & 1438 (886) & 1586 (888) \\\\   NGC 7674\\tablenotemark{2,3} & 0.04 & $<$0.11 & 0.67$^{+0.07}_{-0.05}$ & 0.62$^{+0.47}_{-0.43}$ & 89.0$^{+69.0}_{-40.2}$ & 66.6 (70) & 113 (72) \\\\  \\enddata  \\tablenotetext{1}{All abundances frozen to solar except for Arp 220 which has an abundance of 0.17$^{+0.12}_{-0.05}$.} \\tablenotetext{2}{Best fit absorption same as Galactic absorption. This parameter was then frozen to the Galactic value.} \\tablenotetext{3}{Used c-stat.} \\tablenotetext{4}{ObsIDs 0201830201, 0201830301 and 0201830401. APEC and first power law normalizations fit independently.} \\tablenotetext{5}{ObsID 0013140101.} \\tablenotetext{6}{ObsID 0201830501.} \\tablenotetext{7}{ObsID 0013140201.} \\tablenotetext{8}{ObsID 0554170201.} \\tablenotetext{9}{ObsID 0554170101.} \\tablenotetext{10}{Used pile-up model.} \\end{deluxetable}  \\clearpage \\end{landscape}   \\begin{deluxetable}{lccc}  \\tablewidth{0pt} \\tablecaption{\\label{o3_flux}[OIII] Sample Fluxes (10$^{-14}$ erg cm$^{-2}$ s$^{-1}$)} \\tablehead{  \\colhead{Galaxy} & \\colhead{F$_{0.5-2keV}$} & \\colhead{F$_{APEC}$} & \\colhead{F$_{pow}$} }  \\startdata  Mrk 0609                 &  79.1$_{-3.62}^{+3.62}$ & 4.82$_{-2.15}^{+2.14}$  &  74.3$_{-2.92}^{+2.92}$   \\\\ IC 0486                  &  35.0$_{-8.84}^{+17.6}$ & 0.45$_{-0.39}^{+1.12}$  &  34.6$_{-8.83}^{+17.6}$   \\\\ 2MASX J08035923+2345201  &  1.24$_{-0.71}^{+0.86}$ & 0.83$_{-0.70}^{+0.84}$  &  0.41$_{-0.14}^{+0.17}$   \\\\ 2MASX J08244333+2959238  &  5.25$_{-1.58}^{+1.26}$ & 3.04$_{-1.47}^{+1.04}$  &  2.21$_{-0.59}^{+0.71}$   \\\\ 2MASX J10181928+3722419  &  2.27$_{-0.62}^{+0.92}$ & 1.04$_{-0.43}^{+0.80}$  &  1.23$_{-0.44}^{+0.46}$   \\\\ 2MASX J11110693+0228477  &  1.44$_{-0.59}^{+419}$  & 0.42$_{-0.36}^{+419}$   &  1.02$_{-0.46}^{+1.31}$   \\\\ CGCG 242-028             &  1.17$_{-0.75}^{+2.52}$ & 0.80$_{-0.73}^{+2.50}$  &  0.37$_{-0.19}^{+0.33}$   \\\\ SBS 1133+572             &  3.46$_{-0.82}^{+1.27}$ & 0.52$_{-0.35}^{+0.36}$  &  2.94$_{-0.74}^{+1.22}$   \\\\ Mrk 1457                 &  2.90$_{-1.82}^{+10.3}$ & 2.36$_{-1.79}^{+10.3}$  &  0.54$_{-0.32}^{+0.66}$   \\\\ 2MASX J11570483+5249036  &  3.83$_{-0.66}^{+1.31}$ & 0.91$_{-0.41}^{+1.09}$  &  2.92$_{-0.51}^{+0.72}$   \\\\ 2MASX J12183945+4706275  &  1.54$_{-0.54}^{+2.15}$ & 0.55$_{-0.31}^{+2.10}$  &  0.99$_{-0.44}^{+0.47}$   \\\\ 2MASX J12384342+0927362\t &  6.23$_{-0.76}^{+2.81}$ & 0.86$_{-0.39}^{+2.58}$  &  5.37$_{-0.65}^{+1.12}$   \\\\ NGC 5695                 &  2.57$_{-0.90}^{+5.37}$ & 0.64$_{-0.56}^{+4.31}$  &  1.93$_{-0.71}^{+3.20}$   \\\\  \\enddata \\end{deluxetable}   \\begin{landscape}  \\begin{deluxetable}{lccccl}  \\tablewidth{0pt} \\tablecaption{\\label{12m_flux}12$\\mu$m Sample Fluxes (10$^{-14}$ erg cm$^{-2}$ s$^{-1}$)} \\tablehead{  \\colhead{Galaxy} & \\colhead{F$_{0.5-2keV}$} & \\colhead{F$_{0.5-2keV,extended}$\\tablenotemark{1}} &  \\colhead{F$_{APEC}$} & \\colhead{F$_{pow}$} & \\colhead{Comments}}  \\startdata  NGC 0424        & 31.5$_{-8.53}^{+12.5}$ &  ...                    & 10.6$_{-7.81}^{+12.1}$ &  20.9$_{-3.43}^{+3.26}$ \\\\ NGC 1144        & 8.39$_{-1.61}^{+1.25}$ &  ...                    & 2.87$_{-1.39}^{+0.87}$ &  5.52$_{-0.81}^{+0.90}$ \\\\ NGC 1320        & 28.4$_{-5.46}^{+15.8}$ &  ...                    & 10.5$_{-4.92}^{+15.6}$ &  17.9$_{-2.73}^{+2.38}$ \\\\ NGC 1386        & 25.2$_{-3.99}^{+6.35}$ &  8.48$^{+1.41}_{-1.36}$ & 11.7$_{-3.83}^{+6.19}$ &  13.5$_{-1.12}^{+1.43}$ \\\\ NGC 1667        & 7.31$_{-1.52}^{+1.45}$ &  ...                    & 3.22$_{-1.10}^{+1.03}$ &  4.09$_{-1.05}^{+1.01}$ \\\\ F05189-2524     & 12.1$_{-1.00}^{+2.02}$ &  ...                    & 2.22$_{-0.68}^{+1.88}$ &  9.88$_{-0.74}^{+0.74}$ \\\\ NGC 3982        & 9.14$_{-0.96}^{+0.90}$ &  7.47$^{+1.96}_{-2.49}$ & 3.39$_{-0.73}^{+0.63}$ &  5.75$_{-0.62}^{+0.64}$ \\\\ NGC 4388        & 28.5$_{-7.0}^{+16.0}$  &  23.3$^{+6.6}_{-6.4}$   & 17.8$_{-6.6}^{+15.4}$  &  10.7$_{-2.18}^{+3.15}$ & {\\it Chandra} \\& {\\it XMM-Newton} observations averaged\\\\ NGC 4501        & 7.24$_{-1.30}^{+1.57}$ &  5.47$^{+1.42}_{-1.47}$ & 3.83$_{-0.95}^{+1.06}$ &  3.41$_{-0.88}^{+1.15}$ \\\\ TOLOLO 1238-364 & 14.2$_{4.2}^{+49}$     &  ...                    & 5.43$_{-3.80}^{+5.01}$ &  8.77$_{-1.82}^{+49}$   \\\\ NGC 4968        & 5.85$_{-1.05}^{+1.14}$ &  ...                    & 1.27$_{-0.60}^{+0.67}$ &  4.58$_{-0.86}^{+0.93}$ \\\\ M-3-34-64\t& 50.8$_{-7.81}^{+6.37}$ &  ...                    & 33.4$_{-7.29}^{+4.72}$ &  17.4$_{-2.80}^{+4.28}$ \\\\ NGC 5135        & 34.5$_{-1.68}^{+1.66}$ &  16.4$^{+7.1}_{-5.9}$   & 17.0$_{-1.27}^{+1.23}$ &  17.5$_{-1.10}^{+1.12}$ \\\\ NGC 5194        & 69.6$_{-6.07}^{+6.82}$ &  56.1$^{+8.0}_{-7.9}$   & 52.7$_{-5.97}^{+6.71}$ &  16.9$_{-1.09}^{+1.19}$ \\\\ NGC 5347\t& 2.94$_{-0.46}^{+0.87}$ &  ...                    & 0.50$_{-0.30}^{+0.80}$ &  2.44$_{-0.35}^{+0.34}$ \\\\ Mrk 463         & 14.3$_{-2.72}^{+2.68}$ &  6.04$^{+9.40}_{-1.88}$ & 7.50$_{-2.62}^{+2.59}$ &  6.80$_{-0.73}^{+0.70}$ \\\\ NGC 5506        & 308$_{-53}^{+56}$      &  ...                    & 6.07$_{-4.55}^{+4.48}$ &  302$_{-53}^{+56}$      & ObsIDs 0013140101, 0201830201, 0201830301, 0201830401 \\\\ NGC 5506\t& 444$_{-37}^{+37}$      &  ...                    & 0  \t   \t    &  444$_{-37}^{+37}$      & ObsIDs 0013140201 \\& 0201830501 \\\\ NGC 5506\t& 457$_{-43}^{+37}$      &  ...                    & 10.5$_{-1.52}^{+1.65}$ &  446$_{-43}^{+37}$      & ObsIDs 0554170201 \\& 0554170101 \\\\ Arp 220         & 4.90$_{-1.34}^{+1.19}$ &  4.38$^{+1.11}_{-0.77}$ & 3.05$_{-1.24}^{+0.95}$ &  1.85$_{-0.51}^{+0.71}$ \\\\ NGC 6890        & 11.5$_{-3.76}^{+6.19}$ &  ...                    & 2.62$_{-1.40}^{+1.64}$ &  8.88$_{-3.49}^{+5.97}$ \\\\ IC 5063\t\t& 22.9$_{-3.18}^{+3.12}$ &  4.79$^{+1.15}_{-1.20}$ & 4.57$_{-1.58}^{+1.58}$ &  18.3$_{-2.76}^{+2.69}$ \\\\ NGC7130         & 23.6$_{-1.89}^{+2.63}$ &  8.55$^{+3.65}_{-1.20}$ & 10.3$_{-0.96}^{+1.24}$ &  13.3$_{-1.62}^{+2.32}$ \\\\ NGC7172\t\t& 20.3$_{-2.45}^{+2.28}$ &  ...                    & 2.05$_{-0.51}^{+0.50}$ &  18.3$_{-2.40}^{+2.22}$ & ObsID 0414580101 \\\\ NGC7172\t\t& 11.8$_{-0.72}^{+0.75}$ &  ...                    & 2.06$_{-0.28}^{+0.27}$ &  9.74$_{-0.67}^{+0.70}$ & ObsIDs 0147920601 \\& 0202860101 \\\\ NGC7582\t\t& 40.3$_{-3.7}^{+3.2}$   &  30.8$^{+3.2}_{-2.8}$   & 15.2$_{-3.1}^{+2.5}$   &  25.1$_{-2.0}^{+2.0}$   & {\\it Chandra} \\& {\\it XMM-Newton} observations averaged \\\\ NGC7674\t\t& 17.5$_{-4.95}^{+2.91}$ &  ...                    & 13.6$_{-4.74}^{+2.17}$ &  3.90$_{-1.42}^{+1.94}$ \\\\  \\enddata \\tablenotetext{1}{Sy2s observed with {\\it Chandra} where the unresolved emission after removal of the AGN had high enough signal-to-noise for adequate spectral fitting. F$_{0.5-2keV,extended}$ flux corresponds to this extended emission from the host galaxy (i.e. with the AGN removed).} \\end{deluxetable}  \\end{landscape}  \\clearpage \\begin{landscape} \\begin{deluxetable}{lcccccrr} \\footnotesize \\tablewidth{0pt} \\tablecaption{\\label{pt_src_rem}{\\it Chandra} Spectral Fits to Unresolved Emission Only\\tablenotemark{2}} \\tablehead{ \\colhead{Galaxy} & \\colhead{N$_{H,1}$} & \\colhead{kT$_1$} & \\colhead{kT$_2$} & \\colhead{$\\Gamma$} & \\colhead{N$_{H,2}$}   & \\colhead{$\\chi^2$ 2 APECs} &  \\colhead{$\\chi^2$ 1 APEC} \\\\ \\colhead{ }      & \\colhead{10$^{22}$ cm$^{-2}$} & \\colhead{keV} & \\colhead{keV} & & \\colhead{10$^{22}$cm$^{-2}$} & \\colhead{DOF} & \\colhead{DOF}}  \\startdata   NGC 1386 & $<$0.05 & 0.58$^{+0.09}_{-0.13}$ & ... & 3.23$^{+0.38}_{-0.34}$ & 48.9$^{+32.0}_{-12.4}$ & ... & 79.0 (54) \\\\  NGC 3982\\tablenotemark{2} & 0.01 & 0.20$^{+0.06}_{-0.06}$ & ... & 2.04$^{+0.68}_{-0.68}$ & ... & ... & 46.5 (25) \\\\  NGC 4388\\tablenotemark{2} & 0.03 & 0.60$^{+0.04}_{-0.04}$ & 0.16$^{+0.04}_{-0.02}$ & 0.98$^{+0.51}_{-0.50}$ & ... & 85.4 (83) & 123 (84)\\\\  NGC 4501\\tablenotemark{2} & 0.03 & 0.37$^{+0.22}_{-0.07}$ & ... & 2.26$^{+0.44}_{-0.50}$ & ... & ... & 31.2(38) \\\\  NGC 5135 & $<$0.13 & 0.84$^{+0.09}_{-0.09}$ & 0.39$^{+0.11}_{-0.09}$ & 2.24$^{+0.33}_{-0.31}$ & ... & 62.8 (60) & 77.9 (62) \\\\  NGC 5194\\tablenotemark{2} & 0.02 & 0.60$^{+0.01}_{-0.01}$ & 0.18$^{+0.03}_{-0.02}$ & 3.13$^{+0.10}_{-0.09}$ & 7.92$^{+1.38}_{-1.08}$ & 525 (352) & 616 (353) \\\\  Mrk 463 & $<$0.05 & 0.62$^{+0.05}_{-0.07}$ & ... & 2.50$^{+0.75}_{-0.47}$ & 34.6$^{+10.5}_{-7.7}$ & ... & 86 (71) \\\\  Arp 220\\tablenotemark{2} & $<$0.04 & 0.73$^{+0.07}_{-0.12}$ & ... & 2.10$^{+0.44}_{-0.36}$ & ... & ... & 69.3 (57) \\\\  IC 5063\\tablenotemark{2} & 0.06 & 0.34$^{+0.09}_{-0.04}$ & ... & 1.51$^{+0.55}_{-0.59}$ & 28.5$^{+6.2}_{-5.0}$ & ... & 71.5 (57) \\\\  NGC 7130 & 0.05$^{+0.07}_{-0.03}$ & 0.47$^{+0.07}_{-0.10}$ & ... & 2.24$^{+0.37}_{-0.29}$ & ... & ... & 58.2 (50) \\\\  NGC 7582 & 0.03$^{+0.02}_{-0.02}$ & 0.73$^{+0.04}_{-0.05}$ & ... & 2.37$^{+0.23}_{-0.20}$ & 41.8$^{+10.1}_{-8.2}$ & ... & 142 (128) \\\\  \\enddata  \\tablenotetext{1}{All abundances frozen to solar.} \\tablenotetext{2}{Best fit absorption same as Galactic absorption. This parameter was then frozen to the Galactic value.} \\end{deluxetable} \\end{landscape}   \\clearpage   \\begin{deluxetable}{lccc}  \\tablewidth{0pt} \\tablecaption{\\label{o3_lum}[OIII] Sample Star-formation and AGN Luminosities} \\tablehead{  \\colhead{Galaxy} & \\colhead{Log(L$_{x,SF}$)} &  \\colhead{Log(L$_{x,AGN}$)} }  \\startdata  Mrk 0609                 & 41.38$^{+0.54}_{-0.39}$ & 42.28$^{+0.06}_{-0.06}$ \\\\ IC 0486                  & 40.39$^{+0.87}_{-0.42}$ & 41.83$^{+0.27}_{-0.24}$ \\\\ 2MASX J08035923+2345201  & 40.45$^{+0.87}_{-0.33}$ & 39.64$^{+0.33}_{-0.30}$ \\\\ 2MASX J08244333+2959238  & 40.81$^{+0.60}_{-0.33}$ & 40.18$^{+0.30}_{-0.30}$ \\\\ 2MASX J10181928+3722419  & 40.99$^{+0.72}_{-0.36}$ & 40.54$^{+0.27}_{-0.24}$ \\\\ SBS 1133+572             & 40.76$^{+0.78}_{-0.33}$ & 41.17$^{+0.15}_{-0.15}$ \\\\ Mrk 1457                 & 41.74$^{+0.96}_{-0.42}$ & 40.15$^{+0.57}_{-0.48}$ \\\\ 2MASX J11570483+5249036  & 40.67$^{+0.84}_{-0.36}$ & 40.81$^{+0.12}_{-0.12}$ \\\\ 2MASX J12183945+4706275  & 41.68$^{+0.93}_{-0.42}$ & 41.20$^{+0.33}_{-0.30}$ \\\\ 2MASX J12384342+0927362\t & 41.69$^{+0.87}_{-0.42}$ & 41.90$^{+0.15}_{-0.12}$ \\\\ NGC 5695                 & 40.78$^{+0.96}_{-0.81}$ & 39.79$^{+0.42}_{-0.41}$ \\\\  \\enddata \\end{deluxetable}  \\clearpage  \\begin{landscape} \\begin{deluxetable}{lccl}  \\tablewidth{0pt} \\tablecaption{\\label{12m_lum}12$\\mu$m Sample Star-formation and AGN Luminosities} \\tablehead{  \\colhead{Galaxy} & \\colhead{Log(L$_{x,SF}$)} &  \\colhead{Log(L$_{x,AGN}$)} & \\colhead{Comments}}  \\startdata  NGC 0424        & 40.99$^{+0.69}_{-0.54}$ & 40.72$^{+0.21}_{-0.18}$ \\\\ NGC 1144        & 40.78$^{+0.66}_{-0.24}$ & 40.87$^{+0.21}_{-0.21}$ \\\\ NGC 1320        & 40.33$^{+1.05}_{-0.51}$ & 40.42$^{+0.21}_{-0.21}$ \\\\ NGC 1386\\tablenotemark{1} & 39.23$^{+0.07}_{-0.08}$ & 39.52$^{+0.10}_{-0.07}$ \\\\ NGC 1667        & 40.30$^{+0.60}_{-0.21}$ & 40.06$^{+0.24}_{-0.24}$ \\\\ F05189-2524     & 41.20$^{+0.78}_{-0.30}$ & 41.56$^{+0.15}_{-0.12}$ \\\\ NGC 3982\\tablenotemark{1} & 39.42$^{+0.10}_{-0.18}$ & 38.77$^{+0.04}_{-0.05}$ \\\\ NGC 4388\\tablenotemark{1} & 40.52$^{+0.11}_{-0.14}$ & 39.87$^{+0.19}_{-0.12}$  \\\\ NGC 4501\\tablenotemark{1} & 39.89$^{+0.10}_{-0.14}$ & 39.40$^{+0.09}_{-0.09}$  \\\\ TOLOLO 1238-364 & 40.63$^{+0.72}_{-0.57}$ & 40.31$^{+0.48}_{-0.45}$ \\\\ NGC 4968        & 39.67$^{+0.69}_{-0.33}$ & 39.88$^{+0.12}_{-0.12}$ \\\\ M-3-34-64\t& 41.47$^{+0.45}_{-0.21}$ & 40.84$^{+0.60}_{-0.60}$ \\\\ NGC 5135\\tablenotemark{1} & 40.86$^{+0.16}_{-0.19}$ & 40.90$^{+0.02}_{-0.02}$ \\\\ NGC 5194\\tablenotemark{1} & 39.69$^{+0.06}_{-0.07}$ & 39.08$^{+0.04}_{-0.04}$ \\\\ NGC 5347\t& 39.34$^{+0.84}_{-0.48}$ & 39.47$^{+0.15}_{-0.12}$ \\\\ Mrk 463\\tablenotemark{1} & 41.55$^{+0.41}_{-0.16}$ & 41.69$^{+0.07}_{-0.09}$ \\\\ NGC 5506        & 39.88$^{+0.93}_{-0.27}$ & 41.38$^{+0.12}_{-0.09}$  & ObsIDs 0013140101, 0201830201, 0201830301, 0201830401 \\\\ NGC 5506\t& 0                       & 41.53$^{+0.09}_{-0.06}$  & ObsIDs 0013140201 \\& 0201830501 \\\\ NGC 5506\t& 40.06$^{+0.45}_{-0.21}$ & 41.53$^{+0.06}_{-0.06}$  & ObsIDs 0554170201 \\& 0554170101 \\\\ Arp 220\\tablenotemark{1} & 40.51$^{+0.10}_{-0.08}$ & 39.58$^{+0.09}_{-0.14}$  \\\\ NGC 6890        & 39.83$^{+0.69}_{-0.39}$ & 39.94$^{+0.24}_{-0.24}$ \\\\ IC 5063\\tablenotemark{1} & 40.11$^{+0.09}_{-0.13}$ & 40.69$^{+0.06}_{-0.06}$ \\\\ NGC 7130\\tablenotemark{1} & 40.69$^{+0.15}_{-0.07}$ & 40.94$^{+0.05}_{-0.04}$ \\\\ NGC 7172\t& 39.85$^{+0.39}_{-0.36}$ & 40.45$^{+0.06}_{-0.06}$ & ObsID 0414580101 \\\\ NGC 7172\t& 39.67$^{+0.45}_{-0.15}$ & 40.18$^{+0.12}_{-0.12}$ & ObsIDs 0147920601 \\& 0202860101 \\\\ NGC 7582\\tablenotemark{1} & 40.23$^{+0.04}_{-0.04}$ & 39.72$^{+0.03}_{-0.04}$ \\\\ NGC 7674\t& 41.53$^{+0.51}_{-0.21}$ & 40.03$^{+2.70}_{-0.39}$ \\\\  \\enddata \\tablenotetext{1}{Sy2s observed with {\\it Chandra} where the AGN was removed and extended emission was fitted. L$_{x,SF}$ is derived from these flux values for these objects while L$_{x,AGN}$ is simply L$_{0.5-2keV}$-L$_{x,SF}$. For the remaining sources, L$_{x,SF}$ and L$_{x,AGN}$ is derived as described in the text.} \\end{deluxetable}  \\end{landscape} \\clearpage    \\begin{figure}[ht] \\centering \\subfigure[]{\\includegraphics[scale=0.35,angle=90]{lsf_v_agn_rem.eps}} \\subfigure[]{\\includegraphics[scale=0.35,angle=90]{lagn_v_agn_rem.eps}} \\caption[]{\\label{chandra_v_sings} Comparison of using {\\it Chandra} observations to decompose the soft X-ray emission into star formation and AGN components (L$_{x,SF,Chandra}$ and L$_{x,AGN,Chandra}$, respectively) with estimates of starburst and AGN activity calculated via simulations using XSINGS data of normal galaxies and luminosities derived from spectral fitting (i.e., L$_{x,SF,XSINGS}$, L$_{x,AGN,XSINGS}$). The grey shaded regions enclosed by dashed lines indicate the 3$\\sigma$ confidence interval from a Bayesian linear regression fit \\citep{Kelly}. The dotted-dashed lines denote where the two quantities are equal. These two methods of estimating the star formation and AGN contributions to the soft X-ray emission are consistent.} \\end{figure}   \\begin{figure}[ht] \\centering {\\includegraphics[scale=0.5,angle=90]{soft_ne2_o4_ols.eps}} \\caption[]{\\label{softx_ne2_o4} The results of decomposing the soft X-ray flux (0.5-2 keV) into a star formation and AGN component, using the luminosity of the [NeII] line as a proxy for the former and the luminosity of the [OIV] line to parameterize the latter. The constants $\\alpha$ and $\\beta$ were calculated using ordinary least squares multiple linear regression, where we derive $\\alpha$=1.07 and $\\beta$=0.16. The overplotted line indicates where the two quantities are equal. Cyan triangles represent the 12$\\mu$m Sy2s (with the X-ray variable sources connected by a vertical line), blue diamonds denote the subset of 12$\\mu$m sources with high signal-to-noise {\\it Chandra} imaging of the unresolved emission and red squares mark the [OIII] selected sources.} \\end{figure}    \\begin{figure}[ht] \\centering \\subfigure[]{\\includegraphics[scale=0.40,angle=90]{apec_adj_ne2.eps}} \\subfigure[]{\\includegraphics[scale=0.40,angle=90]{pow_adj_o4.eps}} \\caption[]{\\label{lx_ir}Left: L$_{x,SF}$ vs $\\alpha \\times$L$_{[NeII]}$. Right: L$_{x,AGN}$ vs $\\beta \\times$L$_{[OIV]}$. In both plots, the dashed-dotted line indicates the line of equality, while the gray shaded regions illustrate the 3$\\sigma$ confidence interval from Bayesian linear regression. Color coding is the same as Figure \\ref{softx_ne2_o4}. L$_{x,SF}$ and L$_{x,AGN}$ are derived from spectral fitting of the unresolved {\\it Chandra} emission for the blue diamond data points while these parameters are estimated by scaling L$_{APEC}$ and L$_{pow}$ by as noted in the text. The soft X-ray and IR decomposition approximately agree at the 3$\\sigma$ level, with small deviations appearing at luminosities above 10$^{41}$ erg/s/cm$^2$ for the star formation parameterization and below 2$\\times 10^{41}$ erg/s/cm$^2$ for the AGN decomposition.} \\end{figure}  \\begin{figure}[ht] \\centering {\\includegraphics[scale=0.5,angle=90]{xsfr_v_firsfr.eps}} \\caption[]{\\label{firsfr_xsfr} X-ray derived SFRs (SFR$_{0.5-2keV}$) using our estimate of L$_{x,SF}$ and the calibration of \\citet{PS} as a function of FIR derived SFRs (SFR$_{FIR}$), using {\\it IRAS} 60$\\mu$m and 100$\\mu$m luminosities and \\citet{Kennicutt}'s calibration, for the 12$\\mu$m sample. The overplotted dashed line shows where the two SFRs are equal. A relatively good agreement exists, with the two most significant outliers with the largest SFR$_{FIR}$ values (Arp 220 and F05189-2524) being ULIRGs.} \\end{figure}  \\clearpage"
    },
    "1208/1208.1625_arXiv.txt": {
        "abstract": "{eROSITA will perform the most sensitive X-ray all-sky ever in the energy range $0.3 - 10$ keV. It will likely uncover several $10^4$ compact binaries, most of them  will be cataclysmic variable stars. After a brief introduction to eROSITA we will discuss  the source content of the eROSITA surveys, the expected number of CVs and possible strategies  for the optical follow-up. ",
        "introduction": "eROSITA (extended ROentgen Survey with an Imaging Telescope Array) is an X-ray telescope array built by a German consortium, to be mounted on a Fregat booster and launched with a Russian Zenit rocket  into an L2 halo orbit of the Sun-Earth system in late 2013 \\citep{predehl+10, pavlinsky+09}.  During the first four years of its mission, the eROSITA telescopes will scan the sky in great circles with a scanning speed corresponding to one full circle being completed every four hours. The scanning axis is either pointed  directly towards the sun or a few degrees away from it. As the satellite  moves around the sun, the plane of the scans is thus advanced by about one degree  per day, resulting in a full coverage of the sky every half year.  Hence eROSITA will perform eight independent all-sky surveys in the energy range $0.3-10$\\,keV each lasting half a year. After four years the sky area around the poles of the ecliptic will be exposed for about 20 ks, the average exposure per survey is 240\\,s (eRASS:1) and for the full four year survey about 2000\\,s (eRASS:8). The average sensitivity for the detection of point sources in eRASS:8 will be $1.5 \\times 10^{-14}$\\,\\fcgs\\ and $2\\times 10^{-13}$\\,\\fcgs\\ in the soft ($0.5 - 2.0$\\,keV) and hard ($2 - 10$\\,keV)  energy bands, respectively. Around  the poles of the ecliptic the limiting fluxes will be $4\\times10^{-15}$\\,\\fcgs\\ (soft) and $5\\times10^{-14}$\\,\\fcgs (hard).  The X-ray optics will have an on-axis angular resolution of 15\\arcsec\\ (HEW, half energy width), due to off-axis blurring the average survey resolution will be lower and is expected to be below 30 arc sec. Individual source positions will have a remaining positional inaccuracy of estimated $2-3$ arc sec.  Each of the seven mirror modules has its own camera in its focus, each equipped with a CCD-module and a  processing electronics. The eROSITA-CCDs have $384\\times 384$ pixels which correspond to an image area of 28.8 mm $\\times$ 28.8 mm, respectively, for a field of view of 1.03\\degr\\ diameter. The 384 channels are read out in parallel and the nominal integration time per CCD frame will be 50 ms.  Given the technical specifications, eROSITA can be very much compared to XMM-Newton with EPIC pn, it has however a factor 5 times larger grasp between 0.3 and 2 keV (effective area times FoV), making it a prime survey instrument. During its four-year all-sky survey it shall discover about $10^5$ clusters of galaxies to study the growth of structure and constrain the parameters of Dark Energy as the prime mission goal. Based on the known number-flux relations of AGN \\citep[e.g.][]{gilli}) one may expect the detection of about $3\\times 10^6$\\,AGN of all kind up to redshift 6. The eROSITA X-ray sky will also be populated by about 300,000 coronal emitters. The huge stellar samples built from the large number of detections will help to disentangle the X-ray emitting populations and determine the shape of age-metalicity-activity relations.  Finally, compact galactic (and extragalactic) objects will be detected in large number and in the following sections the expectations are outlined. ",
        "conclusions": ""
    },
    "1208/1208.4000_arXiv.txt": {
        "abstract": "{} {We carried out a radial-velocity survey to search for planets around metal-poor stars. In this paper we report the discovery of two planets around \\mbox{HIP 11952}, a metal-poor star with [Fe/H]$=-1.9$ that belongs to our target sample.} {Radial velocity variations of \\mbox{HIP 11952} were monitored systematically with FEROS at the 2.2~m telescope located at the ESO La Silla observatory from August 2009 until January 2011. We used a cross-correlation technique to measure the stellar radial velocities (RV).} {We detected a long-period RV variation of 290 d and a short-period one of 6.95 d. The spectroscopic analysis of the stellar activity reveals a stellar rotation period of 4.8 d. The Hipparcos photometry data shows intra-day variabilities, which give evidence for stellar pulsations. Based on our analysis, the observed RV variations are most likely caused by the presence of unseen planetary companions. Assuming a primary mass of 0.83 M$_\\odot$, we computed minimum planetary masses of 0.78 M$_\\mathrm{Jup}$ for the inner and 2.93 M$_\\mathrm{Jup}$ for the outer planet. The semi-major axes are $a_1=0.07$ AU and $a_2=0.81$ AU, respectively.} {\\mbox{HIP 11952} is one of very few stars with [Fe/H]$<-1.0$ which have planetary companions. This discovery is important to understand planet formation around metal-poor stars.} ",
        "introduction": "Current results of the exoplanet surveys strongly suggest a correlation between a star's stellar metallicity and its probability of hosting planets, in particular for main-sequence stars (e.g., \\cite{fischer05} 2005; \\cite{johnson10}). According to  these studies, the detection rate of planets decreases with metallicity. However, the conclusions of \\cite{fischer05} (2005) might be affected by an observational bias, since these authors did not have similar numbers of stars in their survey per bin of metallicity. Therefore, it is crucial to understand if either the high stellar metallicity triggers planet formation or the metal enhancement of stars is caused by the formation of planets.  In the last years, exoplanet surveys tried to bridge this gap, starting to include more metal-poor stars in their samples. \\cite{sozzetti09} conducted a three-year RV survey to look for planets around metal-poor stars down to [Fe/H]$=-2.0$ and found no evidence for short-period giant planets within 2~AU from the central star. \\cite{santos11} performed a similar survey, only down to [Fe/H]$=-1.4$ and found three moderately metal-poor stars hosting a long period giant planets (P$>$ 700 d). A hot Saturn and a hot Jupiter have been found transiting two moderate metal-poor stars, with [Fe/H]$=-0.46$ and $-0.4$, respectively (\\cite{bouchy10}; \\cite{simpson11}).  In June 2009 we started a survey to search for planets around metal-poor stars. The target sample includes 96 metal-poor A and F stars. Our target list includes stars with metallicities in the range $-4.0\\le$[Fe/H]$\\le -0.5$. As part of this work, \\cite{setiawan10} found a planet around an extremely metal-poor red horizontal branch star with a short period of 16.2 d. We notice that its [Fe/H]$=-2.1$ is not included in the metallicity range covered by the surveys of \\cite{sozzetti09} and \\cite{santos11}.  \\begin{figure}[t] \\centering \\includegraphics[height=9.0cm,width=11cm,angle=90]{spectraltype.ps} \\caption{Comparison of the spectrum of \\mbox{HIP 11952} with stellar spectra of the Indo-US library. The dashed lines highlight some iron lines (\\ion{Fe}{ii}, \\ion{Ti}{ii}) which are sensitive to the luminosity class of the star. The solid lines in the upper part of the figure indicate the position of some iron and calcium lines which do not vary with the luminosity class.} \\label{spt} \\end{figure}  These recent observations have started to disclose the realm of planets at rather low stellar metallicities, indicating that metallicity may not be the main driver of planet formation. Clearly, more statistics is needed to obtain robust conclusions. In this framework, we report the detection of two planetary companions around \\mbox{HIP 11952} as a result of our RV planet survey around metal-poor stars.  The paper is organized as follows: Observations and data reduction are presented in section~\\ref{obs}. The stellar parameters of \\mbox{HIP 11952} are shown in section~\\ref{star} together with the most relevant information available for this star in the literature. In section~\\ref{rvobs} we describe the RV and photometric analysis. The detection of the planetary companion is addressed in section~\\ref{planet}. Discussion and conclusions are given in sections \\ref{discussion} and \\ref{con}, respectively. ",
        "conclusions": "\\label{con} We observed RV variations of \\mbox{HIP 11952}. The spectroscopic and photometric analysis of the star show that the periodic RV variations are not caused by the intrinsic stellar variability. Based on our analysis, we detected two planetary companions around the metal poor star \\mbox{HIP 11952} with orbital periods of 6.95 d and 290 d. We found evidence for intra-day stellar pulsations and observed a stellar rotation of 4.82 d. We computed the companion's minimum mass of $m_{2} \\sin{i}=0.78$ M$_\\mathrm{Jup}$ for the inner planet and $m_{2} \\sin{i}=2.93$ M$_\\mathrm{Jup}$ for the outer one. Additional high-precise RV measurements are necessary to improve the orbital solution and put more constraints on the eccentricities. Further RV observations might also reveal the presence of other low-mass companions in the system. From the metal abundance analysis that we carried out, we obtained an average [Fe/H]$=-1.90\\pm0.06$ from \\ion{Fe}{i} and \\ion{Fe}{ii}, respectively, which makes \\mbox{HIP 11952} b and c the first planets discovered around a dwarf or subgiant star with [Fe/H]$<$-1.5. This discovery is also remarkable since the planetary system most likely belong to the first generation of planetary systems in the Milky Way."
    },
    "1208/1208.5837_arXiv.txt": {
        "abstract": "{ We investigate  the excitation of magnetoacoustic-gravity waves generated from localized pulses in the gas pressure as well as in vertical component of velocity. These pulses are initially launched at the top of the solar photosphere that is permeated by a weak magnetic field. We investigate three different configurations of the background magnetic field lines: horizontal, vertical and oblique to the gravitational force. We numerically model magnetoacoustic-gravity waves by implementing a realistic (VAL-C) model of solar temperature. We solve two-dimensional ideal magnetohydrodynamic equations numerically with the use of {\\bf{the}} FLASH code to simulate the dynamics of {\\bf{the}} lower solar atmosphere. The initial pulses result in shocks at higher altitudes. Our numerical simulations reveal that a small-amplitude initial pulse can produce magnetoacoustic-gravity waves, which are later reflected from the transition region due to the large temperature gradient. The atmospheric cavities in the lower solar atmosphere are found to be the ideal places that may act as a resonator for various oscillations, including their trapping and leakage into the higher atmosphere. Our numerical simulations successfully model the excitation of such wave modes, their reflection and trapping, as well as the associated plasma dynamics. } ",
        "introduction": "The complicated magnetic field configuration of the Sun plays a key role in various types of dynamical plasma processes in its atmosphere, including all the significant plasma dynamics of the lower solar atmosphere. The resulting  magnetic structures channel energy from the photosphere into the upper atmosphere, in form of magnetohydrodynamic (MHD) waves, and such waves  experience mode conversion, resonances, trapping and reflection which results in the complicated dynamical processes in the lower solar atmosphere, the details of which depend on the plasma properties as well as strength of the magnetic field. The complex magnetic field and plasma structuring in the lower solar atmosphere support the excitation of various kinds of MHD waves, and their propagation, reflection and trapping has been studied extensively both on theoretical and observational grounds (e.g. Murawski et al., 2011; Srivastava \\& Dwivedi, 2010; Srivastava, 2010; Fedun et al., 2009; Srivastava et al., 2008; Hasan et al., 2005; McAteer et al., 2003; Gruszecki et al., 2011, and references therein). The evolving magnetic fields of the lower solar atmosphere also lead to transient processes across a wide range of spatial-temporal scales in form of the eruption and associated phenomena. For example, various types of solar jets are formed at short spatial-temporal scales which play a  significant role in mass and energy transport and also  couple the various layers of the solar atmosphere (e.g. Shibata et al., 2007; Katsukawa et al., 2007; Srivastava \\& Murawski, 2011, and references therein). In addition, the magnetic activity and injections of helicity into the lower solar atmosphere result in large-scale eruptive phenomena, including solar flares and coronal mass ejections (CMEs) in the  outer part of the magnetized solar atmosphere (e.g. Srivastava et al., 2010; Shibata \\& Magara, 2011; Zhang et al., 2012, and references therein). Therefore, the coupling of the complex magnetic field in various layers of the Sun due to waves and transients is one of the most significant areas of contemporary solar research.  In the quiet-Sun magnetic networks, cavities are important locations where the magnetic fields are sufficiently inclined due to their well evolved horizontal components. They are formed over the granular cells in form of the field-free regions due to the transport of plasma at their boundaries and overlaid by bipolar magnetic canopies. The vertical magnetic fields, however, reside mostly in the core of such magnetic networks (Schrijver \\& Title, 2003; Centeno et al., 2007). % Such magnetic cavity-canopy systems are thought to be ideal resonators of the various MHD waves that can be trapped in the cavity, as well can also leak in form of the magnetoacoustic-gravity waves upward through the core of such magnetic networks. It is thought that the field-free cavity regions underlying the bipolar canopy can trap the high-frequency acoustic oscillations, and the low-frequency components may leak into the higher atmosphere in form of magnetoacoustic-gravity waves (Kuridze et al., 2008; Srivastava et al., 2008; Srivastava, 2010). % Therefore, such magnetic structures in the lower solar atmosphere are the regions that may play an important role in the wave filtering (e.g. McIntosh \\& Judge, 2001; Krijger et al., 2001; McAteer et al., 2002; Vecchio et al., 2007, 2009; Srivastava, 2010, and references therein).  Alongside recent high-resolution observations of MHD waves in the lower solar atmosphere, extensive efforts have been {\\bf{made}} in the area of analytical and numerical modeling of such waves: Fedun et al. (2009) have {\\bf{investigated}} the 3D numerical modeling of the coupled slow and fast magnetoacoustic wave propagation in the lower solar atmosphere; and recently Fedun et al. (2011a) have {\\bf{reported the first}} numerical results of the frequency filtering of torsional Alfv\\'en waves in the chromosphere. Apart from general numerical modeling of the waves in the lower solar atmosphere, models have also {\\bf{investigated}} the acoustic wave spectrum in the localized magnetic structures of the lower solar atmosphere, e.g., magnetic cavity-canopy system (e.g. Kuridze et al., 2008; Srivastava et al., 2008; Kuridze et al., 2009, and references cited therein).  It is also noteworthy that Bogdan et al. (2003), Fedun et al. (2009) and Fedun et al. (2011a) discussed in detail the excitation, propagation and conversion of magnetoacoustic waves in a realistic 3D MHD simulation. However, in these references, the waves driven by a periodic driver were discussed, whereas we numerically simulate the excitation of magnetoacoustic-gra\\-vi\\-ty waves generated by pulses in the gas pressure and vertical component of velocity, mimicking an isolated solar granule. We aim to investigate and understand this simpler (albeit complex enough) system before tackling the more realistic, multiple granule system. Our philosophy is to build up our models incrementally, with a clear focus on the underlying physical processes at each step.  In this paper, we  investigate  the excitation of magnetoacoustic-gravity waves generated from localized pulses in the gas pressure as well as in vertical component of velocity, modelling the effect of an isolated solar granule. These pulses are initially launched at the top of the solar photosphere that is permeated by a weak magnetic field. We investigate three different configurations of the background magnetic field lines: vertical, horizontal, and oblique to the gravitational force. We aim to show that small amplitude perturbations that are associated with such a granule are able to trigger large amplitude, complicated oscillations in the solar corona, which exhibit periodicities within the detected range of $3-5$ min.  The structure of the paper is as follows: In Sect.~\\ref{SECT:NUM_MODEL} we describe the numerical model. We report the numerical results in Sect.~\\ref{sect:results} and present the discussion and conclusions in {\\bf{Sect.~\\ref{SECT:DISS} }}. ",
        "conclusions": "\\label{SECT:DISS} We have investigated the impulsive excitation of magnetoacoustic-gravity oscillations and have compared and contrasted their resultant propagation under different orientations of equilibrium magnetic fields. We performed 2D numerical simulation of the velocity and gas pressure pulses, which mimic a solar granule. These pulses were initially launched at the top of solar photosphere, in a stratified solar atmosphere utilizing the VAL-C temperature profile. We find that in the cases the background magnetic field possesses a non-zero vertical component the amplitude of the upwardly-propagating perturbations rapidly grows with height due to the rapid decrease in the equilibrium mass density. Therefore, the perturbation quickly steepens into shocks in the upper regions of the solar chromosphere that launches the cool material behind it.   We find that the excitation of magnetoacoustic-gravity waves in a non-ho\\-ri\\-zon\\-tal equilibrium magnetic field configuration is channeled upwards and is able to cause large amplitude transition region oscillations. Meanwhile, the horizontal equilibrium magnetic field configuration results in comparatively smaller amplitude transition region oscillations, due to fast magnetoacoustic-gravity waves which spread their energy across magnetic field lines in contrast to slow magnetoacoustic-gravity waves which in a strongly magnetized plasma region are guided along magnetic field lines. This indicates that the wave energy is transfered into the upper atmosphere in larger amounts where the magnetic field is more vertical. However, the complexity and therefore the evolution of a horizontal equilibrium magnetic field allows the reflection and trapping of the waves in the lower solar atmosphere and can influence the localized dynamics and heating of the atmosphere. The reflection of the waves in the lower solar atmosphere and its trapping may also result in the presence of chromospheric cavities.  By analyzing the time-signature of $V_{\\rm{y}}$ collected at a fixed spatial point, we also find that the period of oscillation is {\\emph{longer}} when comparing the vertical magnetic field case (here, the characteristic period of oscillation was $\\approx 250-300$ s) with the horizontal equilibrium magnetic system (in which the characteristic period of oscillation was $\\approx 150-200$ s). For the vertical magnetic field system in low plasma $\\beta$ regions, the magnetoacoustic-gravity waves are well-described as {\\emph{slow}} magnetoacoustic-gravity waves (propagating along the magnetic field lines at approximately the sound speed, $c_{\\rm s}(y)$, see Eq.~\\ref{eq:c_s}). For the horizontal magnetic field system in the strongly magnetized plasma, the oscillations transverse to the magnetic field are well-described as {\\emph{fast}} magnetoacoustic-gravity waves (propagating across the magnetic field lines at the fast speed, $c^2_{\\rm f}(y) = c_{\\rm A}^2(y) + c_{\\rm s}^2(y)$, see Eqs.~\\ref{eq:c_A} and \\ref{eq:c_s}). As for $y=0.5$ Mm we have $c_{\\rm A}\\ll c_{\\rm s}$, therefore the cut-off frequencies of the fast and slow magnetoacoustic-gravity waves are very close to each other, which results in very close values of the cut-off waveperiods. Therefore, we could expect that the detected waveperiods for the vertical and horizontal magnetic fields exhibit similar values. However, the cut-off waveperiods are derived on the basis of the linear theory which is valid for small amplitude oscillations only. As such small oscillations are present in the case of the horizontal magnetic field, therefore, the obtained numerical data lies close to the analytical prediction. In the case of the vertical magnetic field, the upwardly propagating waves interact with the waves which become reflected from the transition region. As a result of that larger amplitude oscillations originate, which significantly alter the background plasma. The waves, which propagate through this strongly modified medium exhibit modified velocities and they get reflected from the transition region that is locally largely curved. As a consequence of that waveperiods within the range $250-300$ s result, which differ from $P_{\\rm ac}$.   {\\bf It is known that radiation is an effective mechanism of wave damping in the low photosphere ((Mihalas \\& Toomre, 1982).% It might not radically alter the system's behavior, but radiative damping is at least likely to reduce the amplitude of the waves reaching the transition region, leading to shorter jets than what we see in our models, as well as lower amplitudes of the coronal shocks. In this paper, we do not invoke the radiative cooling nor thermal conduction in our model atmosphere as we aim to model the small-scale atmospheric regions above the solar photosphere where such effects are not believed to be dominant. However, we intend to include these effects in our future studies.  It is noteworthy that the solar atmosphere is structured by convective overshoot which is absent in the model we devised. Instead, we isolated a single granule-like perturbation by aiming to mimic a magnitude of flow and plasma temperature that are associated with the solar granulation (Baran 2011), and considered the complex scenario which results by this simple model to understand the physics of wave phenomena above such magnetic structures in the solar atmosphere. We intend to develop more advanced models in our future studies.}  In conclusion, our numerical simulations clearly demonstrate that small amplitude initial pulses in vertical velocity and gas pressure are able to trigger a plethora of dynamic phenomena in the upper regions of the solar atmosphere with waveperiods within the range of $150-300$ s, a value which depends on orientation of the background magnetic field. However, it should be noted that the performed 2D simulations are idealized in the sense that they do not include radiative transfer and thermal conduction along field lines. The magnetic field configuration and the equilibrium stratification are simple and we modeled a single granule only. These limitations require additional studies which we intend to carry on in near future.  \\begin{acks} {\\bf The authors express their thanks to the referee for their stimulating comments. } This work has been supported by a Marie Curie International Research Staff Exchange Scheme Fellowship within the 7th European Community Framework Program (K.M.). This research was carried out with the support of the \"HPC Infrastructure for Grand Challenges of Science and Engineering\" Project, co-financed by the European Regional Development Fund under the Innovative Economy Operational Program (K.M.). The software used in this work was in part developed by the DOE-supported ASC/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago. JM acknowledges IDL support provided by STFC, UK. AKS thanks Shobhna Srivastava for encouragements and supports. RO acknowledges financial support from MICINN/MINECO and FEDER funds through grant AYA2011-22846 and also CAIB through the ``Grups Competitius'' scheme and FEDER funds. KM expresses his thanks to Kamil Murawski for his assistance in drawing the numerical data. \\end{acks}"
    },
    "1208/1208.2542_arXiv.txt": {
        "abstract": "The current concordance model of cosmology is dominated by two mysterious ingredients: dark matter and dark energy. In this paper, we explore the possibility that, in fact, there exist two dark-energy components: the cosmological constant $\\Lambda$, with equation-of-state parameter $w_\\Lambda=-1$, and a `missing matter' component $X$ with $w_X=-2/3$, which we introduce here to allow the Friedmann equation written in terms of conformal time $\\eta$ to be form-invariant under the reciprocity transformation $a(\\eta) \\to 1/a(\\eta)$ of the scale factor.  Using recent cosmological observations, we constrain the present-day energy density of missing matter to be $\\Omega_{X,0}=-0.11\\pm 0.14$. This is consistent with the standard $\\Lambda$CDM model, but constraints on the energy densities of all the components are considerably broadened by the introduction of missing matter; significant relative probability exists even for $\\Omega_{X,0} \\sim 0.2$, and so the presence of a missing matter component cannot be ruled out. Nonetheless, a Bayesian model selection analysis disfavours its introduction by about 1.5 log-units of evidence. Foregoing our requirement of form invariance of the Friedmann equation under the reciprocity transformation, we extend our analysis by allowing $w_X$ to be a free parameter. For this more generic `double dark energy' model, we find $w_X = -1.02 \\pm 0.20$ and $\\Omega_{X,0} = 0.08 \\pm 0.57$, which is again consistent with the standard $\\Lambda$CDM model, although once more the posterior distributions are sufficiently broad that the existence of a second dark-energy component cannot be ruled out. Moreover, the two-dimensional posterior in the $(w_X,\\Omega_{X,0})$-plane is strongly bimodal with both peaks offset from the standard $\\Lambda$CDM model corresponding to $(-1,0)$, although the latter is still admissible; this bimodality is in contrast to the correctly-centred unimodal posterior obtained when analysing simulated observations from a $\\Lambda$CDM model. The model including the second dark energy component also has a similar Bayesian evidence to $\\Lambda$CDM to within the numerical uncertainty. ",
        "introduction": "Over the past decade, cosmological observations have confirmed that the background expansion of the universe is accelerating \\cite{Riess98, Perlmutter99}. This remarkable phenomenon is usually explained by assuming the existence of a single dark-energy component, often modelled as a perfect fluid with a (generally time-dependent) equation-of-state parameter $w(z)$ that results in it exhibiting a negative pressure. The simplest form of dark energy is a cosmological constant $\\Lambda$, which corresponds to a \\textit{constant} equation of state $w_{\\Lambda}=-1$. Together with cold dark matter, which is key to explaining the formation of structure in the universe, the cosmological constant gives rise to the standard $\\Lambda$CDM model, which provides a good fit to existing cosmological observations.  Nonetheless, there have been a large number of other exotic forms of matter proposed to provide alternative explanations for the current accelerating universal expansion \\cite{Copeland06,Durrer08}, including, for example, topological defects \\cite{Vilenkin85}.  In this paper, we remain focussed on the $\\Lambda$CDM model, but with the inclusion of a second dark energy component, which is introduced (in the first instance) to allow the Friedmann equation written in terms of conformal time $\\eta$ to be form invariant under the reciprocity transformation $a(\\eta)\\to 1/a(\\eta)$ of the universal scale factor \\cite{Ibison11}. Such an invariance is of general interest, but may be particularly relevant for Penrose's recent `Cycles of Time' cosmological model \\cite{Penrose}, which posits a cyclic universe in which the ultimate infinitely expanded state of one phase (or `aeon') is identified with the initial singularity of the next.  For a homogeneous and isotropic universe described by the Friedmann--Robertson--Walker (FRW) metric with curvature parameter $k$, the Friedmann equation describing the dynamical evolution of the scale factor $a(t)$ can be written as \\begin{equation}\\label{eq:Fried} \\left( \\frac{H}{H_0} \\right)^2 = \\sum_i \\Omega_{i,0} \\, a^{-3 (1+w_i)}, \\end{equation} \\noindent where $H=\\dot{a}/a$ is the Hubble parameter (the dot denotes differentiation with respect to cosmic time $t$), and the energy density $\\rho_i$ of each of the constituent components of the universe is taken into account through the corresponding present-day density parameter $\\Omega_{i,0} = 8\\pi G \\rho_{i,0}/(3H_0^2)$ and its equation-of-state parameter $w_i$, which we will assume throughout to be time-independent.  In the $\\Lambda$CDM model, the total density parameter is usually taken to comprise of contributions from radiation ($w=\\frac{1}{3}$), matter (typically modelled as dust with $w=0$), curvature ($w=-\\frac{1}{3}$), and the cosmological constant ($w=-1$).  These are listed in Table~\\ref{tab:omega}, in which one can see an obvious `gap' that we term `missing matter' with $w=-\\frac{2}{3}$.  Interestingly, forms of matter do exist for which $w=-\\frac{2}{3}$, such as domain walls \\cite{Conversi04, Battye99,Mithani12}, or particular scalar field models \\cite{Caldwell98}. \\begin{table} \\begin {center} \\begin{tabular}{rcl} \\hline \\vspace{0.1cm} $w_i$\t\t\t\\qquad \\qquad\t \t&   {\\rm component} & \\qquad $\\Omega_i$ \\\\ \\hline \\vspace{0.1cm} 1/3\t\t\t\t\\qquad\t\\qquad\t& {\\rm radiation}\t & \\qquad $\\Omega_{\\rm r}$ \t\\\\ \\vspace{0.1cm} 0\t\t\t\t\\qquad\t\\qquad         & {\\rm matter (dust)}\t\t& \\qquad $\\Omega_{\\rm m}$   \t\\\\ \\vspace{0.1cm} $-1/3$\t\t\t\t\\qquad\t\\qquad\t&{\\rm curvature}\t& \\qquad $\\Omega_{k}$\t\\\\ \\vspace{0.1cm} $-2/3$\t\t\t\t\\qquad\t\\qquad\t& {\\rm missing matter ?} \t& \\qquad $\\Omega_{X}$\t\t\\\\ \\vspace{0.1cm} $-1$\t\t\t\t\\qquad\t\\qquad\t& cosmological constant \t& \\qquad $\\Omega_{\\Lambda}$ \\\\ \\hline \\end{tabular} \\caption{Equation-of-state parameters for different constituents of the universe.} \\label{tab:omega} \\end{center} \\end{table} More importantly, however, the inclusion of a component of precisely this form is \\textit{required} if the Friedmann equation written in terms of conformal time $\\eta$ is to be form invariant under the reciprocity transformation $a(\\eta)\\to 1/a(\\eta)$ \\cite{Ibison11}. This is easily seen by first making the change of variable $d\\eta = dt/a$ in the Friedmann equation (\\ref{eq:Fried}), and including an additional missing matter component $X$; this reads \\begin{equation} \\frac{1}{H_0^2}\\left(\\frac{da}{d\\eta}\\right)^2 = \\Omega_{{\\rm r},0}+\\Omega_{{\\rm m},0}a + \\Omega_{k,0}a^2 + \\Omega_{X,0} a^3 + \\Omega_{\\Lambda,0} a^4. \\label{eq:friedfull} \\end{equation} Making the change of variable $\\tilde{a}(\\eta)=1/a(\\eta)$, one then immediately finds that \\begin{equation} \\frac{1}{H_0^2}\\left(\\frac{d\\tilde{a}}{d\\eta}\\right)^2 = \\Omega_{{\\rm r},0}\\tilde{a}^{4}+\\Omega_{{\\rm m},0}\\tilde{a}^{3} + \\Omega_{k,0}\\tilde{a}^2 + \\Omega_{X,0} \\tilde{a} + \\Omega_{\\Lambda,0}. \\label{eq:friedinvert} \\end{equation} Given the observed values of the various present-day density parameters $\\Omega_{i,0}$, this equation is clearly not identical to (\\ref{eq:friedfull}). Nonetheless, we do find that, provided $\\Omega_{X,0} \\neq 0$, the Friedmann equation is \\textit{form} invariant, with a swapping of roles between radiation and the cosmological constant, and between matter and our additional `missing matter' component (hence our name for it). We note that the curvature density preserves its dynamical role under the reciprocity transformation and also that the remaining Einstein equations are satisfied for the new `$w$' values implied by (\\ref{eq:friedinvert}), showing that the reciprocity transformation is a symmetry of the entire set of cosmological equations \\footnote{It is further shown in \\cite{Ibison11} that, provided $\\Omega_{X,0} \\neq 0$, the Friedmann equation (\\ref{eq:friedfull}) is, in fact, form invariant under the more general M\\\"obius transformation $\\tilde{a}(\\eta)=[\\alpha-\\gamma\\,a(\\eta)]/[\\beta-\\delta\\,a(\\eta)]$, where $\\alpha$, $\\beta$, $\\gamma$ and $\\delta$ are constants, of which the reciprocity transformation $\\tilde{a}(\\eta)=1/a(\\eta)$ is a special case (note that there are only three effective constant degrees of freedom in the M\\\"obius transformation; a convenient choice is to set $\\delta\\alpha-\\beta\\gamma=1$). Under the general M\\\"obius transformation, however, one finds that the Friedmann equation (\\ref{eq:friedfull}) becomes \\begin{equation*} \\frac{1}{H_0^2}\\left(\\frac{d\\tilde{a}}{d\\eta}\\right)^2 = \\sum_{i=0}^4 \\tilde{\\Omega}_{i,0}\\tilde{a}^i, \\end{equation*} where $\\tilde{\\Omega}_{i,0} = \\tilde{\\Omega}_{i,0}(\\alpha,\\beta,\\gamma,\\delta, \\Omega_{{\\rm r},0},\\Omega_{{\\rm m},0},\\Omega_{k,0},\\Omega_{X,0}, \\Omega_{\\Lambda,0})$ are a new set of `density parameters', each of which, in general, depends on \\textit{all} the original density parameters, as well as the parameters in the M\\\"obius transformation. Thus one loses the simple swapping of roles between different original energy densities that occurs for the reciprocity transformation.}. \\\\  Once one admits the possibility of adding an extra component to the energy content of the universe, it is natural to extend one's investigation by allowing its equation-of-state parameter to vary, rather than fixing it to $w=-\\frac{2}{3}$. This more generic `double dark energy' model comes at the cost of breaking the reciprocity invariance of the Friedmann equation, but it is also of interest in its own right since the observed acceleration of the universal expansion may be driven by more than just a single dark-energy component.  The structure of this paper is as follows. In Section~\\ref{sec:phenom}, we give a brief summary of the phenomenology of an additional missing matter component with $w=-\\frac{2}{3}$ by investigating its effect on the expansion history of the universe, in particular the distance-redshift relation, and on the evolution of perturbations, through the cosmic microwave background (CMB) and matter power spectra. In Section~\\ref{sec:analysis}, we describe our Bayesian parameter estimation and model selection analysis methodology and the cosmological data sets used to set constraints on our `missing matter' and `double dark energy' models. The results of these analyses are given in Section~\\ref{sec:results} and our conclusions are presented in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}   We have investigated the possibility that there exist two dark-energy components in the universe: a cosmological constant, with $w=-1$; and an additional component $X$ with equation-of-state parameter $w_X$. In the first instance, we fix the equation-of-state parameter of $X$ to the value $w_X=-\\frac{2}{3}$. This `missing matter' model corresponds to the special case in which the additional component is \\textit{required} for the Friedmann equation written in terms of conformal time $\\eta$ to be form invariant under the reciprocity transformation $a(\\eta) \\to 1/a(\\eta)$. Foregoing this requirement, we then consider the more general `double dark energy' model, in which $w_X$ is a free parameter assumed to have uniform prior in the range $w_X=[-\\frac{3}{2},-\\frac{1}{2}]$. For both models, we perform a Bayesian parameter estimation and model selection analysis, relative to standard $\\Lambda$CDM, using recent cosmological observations of cosmic microwave background anisotropies, Type-IA supernovae and large scale-structure, together with constraints on the baryon density from Big Bang Nucleosynthesis and on the Hubble parameter from the Hubble Space Telescope key project.  For the missing matter model, the introduction of the additional component $X$ significantly broadens the constraints on the standard parameters in the $\\Lambda$CDM model, but leaves their best-fit values largely unchanged. The 1D marginalised constraint on the missing matter density parameter is $\\Omega_{X,0} =-0.11\\pm 0.14$.  Thus, current cosmological observations prefer a slightly negative value, which is difficult to interpret physically, but the posterior on this parameter is sufficiently broad that significant relative probability exits even for $\\Omega_{X,0} \\sim 0.2$, and so the presence of a missing matter component cannot be ruled out. Nonetheless, our results are consistent with $\\Lambda$CDM and our Bayesian model selection analysis disfavours the missing matter model, as compared to $\\Lambda$CDM, by about 1.5 log-units of evidence.  For the double dark energy model, the constraints on standard $\\Lambda$CDM parameters are again considerably broadened. The 1D marginalised constraints on the second dark energy component are $\\Omega_{X,0}= 0.080\\pm 0.574$ and $w_X=-1.02\\pm 0.20$, respectively, which are again consistent with $\\Lambda$CDM. Once more, however, the 1D marginalised posterior on $\\Omega_{X,0}$ is sufficiently broad that even $\\Omega_{X,0} \\sim 1.0$ is not ruled out. More interestingly, the 2D marginal distributions in the $(w_X, \\Omega_{\\Lambda,0})$ and $(w_X,\\Omega_{X,0})$ planes are both bi-modal, exhibiting a `butterfly' shape. In particular, the peaks in the $(w_X,\\Omega_{X,0})$-plane are offset from the $\\Lambda$CDM value $(-1,0)$, although the latter is still acceptable. The two modes of the distribution correspond to models with $\\Omega_{\\Lambda,0} \\approx 0.3$, $\\Omega_{X,0} \\approx 0.4$, $w_X \\approx -1.15$ and $\\Omega_{\\Lambda,0} \\approx 1.05$, $\\Omega_{X,0} \\approx -0.35$, $w_X \\approx -0.88$, respectively. We also find that the double dark energy model has a similar Bayesian evidence to $\\Lambda$CDM to within the numerical uncertainty, and hence neither model is preferred over the other. \\begin{figure*} \\includegraphics[trim = 15mm 60mm 15mm 60mm, clip, width=11cm, height=9cm]{Noise_wx_ok.pdf} \\caption{ 1D and 2D marginalised posterior distributions for density parameters in the double dark energy model (note that $\\Omega_{m,0} = 1-\\Omega_{\\Lambda,0}-\\Omega_{k,0}-\\Omega_{X,0}$) as derived from simulated observations of the CMB power spectrum and Type Ia supernovae generated assuming a concordance $\\Lambda$CDM cosmology. The 2D constraints are plotted with $1\\sigma$ and $2\\sigma$ confidence contours. \\label{fig:XCDMsim}} \\end{figure*} To investigate the significance of the observed bimodality of the posterior, we performed simulated observations of the CMB power spectrum and Type IA supernovae apparent brightness versus redshift, assuming a standard concordance $\\Lambda$CDM model and observational data quality commensurate with the real CMB and supernovae data used in our analysis. These were combined with the same constraints on the baryon density from BBN \\citep{BBN} and Gaussian prior on $H_0$ from the HST key project \\citep{HST} that were used in the analysis of the real data.  The resulting posterior distribution of the density parameters is illustrated in Fig.~\\ref{fig:XCDMsim}. As one might expect, the 2D marginal distributions in the $(w_X, \\Omega_{\\Lambda,0})$ and $(w_X,\\Omega_{X,0})$ planes do exhibit a characteristic `plus' shape, which is indicative of the parameter degeneracies along the coordinate axes.  In contrast to the bimodal posterior obtained from the real data, however, we see that the analysis of simulated observations yields a unimodal posterior centred correctly on the input concordance parameters values. This suggests that the bimodal nature of the posterior in Fig.~\\ref{fig:XCDM} may be driven by features present in the real data, but absent from the $\\Lambda$CDM simulation.   Since the double dark energy model remains viable with current cosmological observations, it is of interest to generalise it further. The bimodality we observe in Figure~\\ref{fig:XCDM} might be evidence of contributions from two species $X$ and $Y$, with $w_X \\approx -1.2$ and $w_Y \\approx -0.8$. It might therefore be of interest to remove the requirement of having one fluid with $w=-1$ and investigate if the value of either $w_X$ or $w_Y$ settles near $-1$. Also of interest would be to allow the equation-of-state parameter $w_X(z)$ of the second dark energy component (and possibly also of $Y$) to depend on redshift \\cite{Vazquez12b}. We plan to investigate these possibilities in a future work."
    },
    "1208/1208.6011_arXiv.txt": {
        "abstract": "{We compute the covariant three-point function near horizon-crossing for a system of slowly-rolling scalar fields during an inflationary epoch, allowing for an arbitrary field-space metric. We show explicitly how to compute its subsequent evolution using a covariantized version of the separate universe or `$\\delta N$' expansion, which must be augmented by terms measuring curvature of the field-space manifold, and give the nonlinear gauge transformation to the comoving curvature perturbation. Nonlinearities induced by the field-space curvature terms are a new and potentially significant source of non-Gaussianity. We show how inflationary models with non-minimal coupling to the spacetime Ricci scalar can be accommodated within this framework. This yields a simple toolkit allowing the bispectrum to be computed in models with non-negligible field-space curvature.} ",
        "introduction": "\\label{sec:introduction}  Considerable effort has recently been invested in the study of multiple-field models of inflation. There are three principal motivations: First, unless curvature couplings flatten their potentials at high energy density, the Standard Model has not produced scalar fields which can successfully inflate. This has stimulated the search for realizations of inflation in theories beyond the Standard Model. Second, single-field inflationary models often require super-Planckian field excursions. Unfortunately, well-rehearsed arguments suggest that we should not expect the scalar potential to be stable against renormalization-group running over such large distances in field space. Multiple-field models may evade this problem by allowing sub-Planckian excursions. Finally, interactions between several fields can give rise to observable non-Gaussianity. This may make some multiple-field models sufficiently predictive that they can be falsified by observation.  It has been known for some time that, under slow-roll conditions, multiple-field models with canonical kinetic terms generate unobservable three- and four-point functions at horizon-crossing \\cite{Seery:2005gb,Seery:2006vu,Seery:2008ax,Chen:2010xka}. Observable effects can arise only later, from nonlinear processes operating on superhorizon scales. A reasonably clear picture has emerged in which these nonlinear effects can be understood as the deformation of a Gaussian probability distribution by the phase space flow associated with the theory \\cite{Mulryne:2009kh,Mulryne:2010rp,Elliston:2011dr,Seery:2012vj}.  Noncanonical kinetic terms offer new possibilities. In some theories the Lagrangian becomes an arbitrary function of the kinetic energy $2X = - \\partial_\\mu \\vp \\partial^\\mu \\vp$. Where this yields a reduced sound speed for perturbations there can be a significant enhancement of the three- and four-point functions \\cite{Alishahiha:2004eh,Seery:2005wm,Chen:2006nt,Cheung:2007st, Arroja:2008yy,Langlois:2008qf,RenauxPetel:2011uk}. But this is not the only type of noncanonical Lagrangian. In many examples descending from our present ideas about physics at very high energies, including string theory and supergravity, the kinetic energy must be written $2X = - \\G_{IJ} \\partial_\\mu \\vp^I \\partial^\\mu \\vp^J$, where $\\G_{IJ}$ is an arbitrary, symmetric function of the fields $\\vp^I$. The simplest example is the nonlinear $\\sigma$-model of Gell-Mann~\\&~L\\'{e}vy, originally introduced to describe spin-0 mesons.  The matrix $\\G_{IJ}$ can be thought of as a metric on the space of scalar fields and will generically exhibit nonzero curvature. Are there interesting enhancements of non-Gaussian effects in these models? Estimates of the three-point function have been made by a number of authors~\\cite{Sun:2005ig,Langlois:2008qf,RenauxPetel:2011uk, Peterson:2011yt,Saffin:2012et}, but as yet no complete formalism exists which allows the bispectrum to be followed from horizon-crossing up to the time of observation. Moreover, Gong \\& Tanaka recently pointed out that the most widely-used formulation of nonlinear perturbation theory is not covariant \\cite{Gong:2011uw}. They introduced a covariant description, to be discussed in \\S\\ref{sec:covariant-perts}, and constructed the action for fluctuations up to third-order in the scenario of Langlois~et~al.~\\cite{Langlois:2008qf}. A similar argument was later made by Saffin~\\cite{Saffin:2012et}. Covariance is a convenience rather than a physical principle, so its absence does not invalidate earlier results. Nevertheless, it is a considerable convenience: there are subtleties associated with time evolution of the two-point function on curved field-space which are most clearly expressed in covariant language. These take the form of a contribution to the effective mass-matrix from the Riemann curvature tensor. In this paper we show that a similar phenomenon occurs for the three-point function.  In flat field-space, time evolution of superhorizon modes may be taken into account using the `separate universe' method~\\cite{Starobinsky:1982ee,Lyth:1984gv,Starobinsky:1986fxa,Lyth:2005fi}. This enables the time dependence of each fluctuation to be determined from the relative behaviour of separated spatial regions following slightly displaced phase-space trajectories. It can be effected using a Taylor expansion to compare two solutions of the slow-roll equation. But in curved field-space we must be cautious when comparing the relative motion of neighbouring trajectories. In the analogous case of general relativity one would use the equation of geodesic deviation, or `Jacobi equation'. In flat field-space this can be integrated to reproduce the Taylor expansion~\\cite{Seery:2012vj}. When promoted to curved field-space the Jacobi approach is automatically covariant and accounts naturally for time-dependent effects generated by the Riemann curvature, including its known contribution to the effective mass-matrix. Beyond linear order there are new contributions which influence the three-point function. These must appear in any correct formulation but are expressed most clearly and economically in terms of the field-space curvature.  At present, the covariant approach cannot be used to generate observable predictions beyond the power spectrum. To do so would require a determination of the initial value of the three-point function near horizon-crossing, together with a prescription to evolve it into the primordial curvature perturbation. It is only the curvature perturbation which can be connected with observable quantities. Neither of these pre-requisites has yet been provided.  In this paper we compute the initial value of the covariant three-point function at horizon exit and use the Jacobi approach to determine its time evolution. (Partial expressions for the noncovariant three-point function were given by Langlois~et~al.~\\cite{Langlois:2008qf} and Renaux-Petel~et~al.~\\cite{RenauxPetel:2011uk}.) As a concrete example we give the analysis for the $\\sigma$-model Lagrangian $\\mathcal{L} = X + V$, although our methods generalize to more complex cases. We show that this initial value can be smoothly connected to the subsequent Jacobi evolution. In particular, the evolutionary effects described above which depend on the Riemann tensor can be matched to new infrared divergences in the three-point function. We demonstrate this matching explicitly to subleading order in both time-dependent perturbation theory and the slow-roll expansion. A further benefit of the Jacobi approach is that time evolution can be computed very simply using ordinary differential equations.  In~\\S\\ref{sec:perturbations} we specialize the results of Gong~\\&~Tanaka to the $\\sigma$-model Lagrangian and obtain the action for fluctuations to third-order. In~\\S\\ref{sec:n-pfs} we compute the corresponding two- and three-point functions near horizon-crossing in the spatially flat gauge. The two-point function has been known since the work of Sasaki~\\&~Stewart~\\cite{Sasaki:1995aw}, but the computation of the covariant three-point function is new. In~\\S\\ref{sec:evolution} we use the Jacobi method discussed above to compute the evolution of these correlation functions after horizon-crossing.  In~\\S\\ref{sec:action} we show that our results can be applied to models in which the scalar fields are coupled non-minimally to gravity by making a suitable conformal transformation \\cite{Kaiser:2010ps}. Such couplings arise naturally in the low-energy limit of higher-dimensional theories including supergravity, string theory and Kaluza--Klein models \\cite{Rainer:1996gw,Gunther:1997ft,Appelquist:1987nr}, or as counterterms in curved spacetime \\cite{Birrell:1982ix,Buchbinder:1992rb}. Finally, we conclude in~\\S\\ref{sec:conclusion}.  \\paragraph{Notation.} Throughout, we work in units where $c = \\hbar = 1$ and express the gravitational coupling in units of the reduced Planck mass, $\\Mpl^{-2} \\equiv 8 \\pi G$. Upper-case Latin indices $I$, $J$, $K$, \\ldots, label the species of scalar fields, and Greek letters label spacetime indices. We use a modified index convention for bilocal tensors (`bitensors'), to be described in \\S\\ref{sec:third_order}. The covariant derivative compatible with the field-space metric $\\G^{IJ}$ is $\\grad_I$. For any tensor $\\mathcal{F}_{\\cdots}$ we write $\\grad_I \\mathcal{F}_{\\cdots} = \\mathcal{F}_{\\cdots;I}$. Our sign convention for the curvature tensor is defined by the Ricci identity, $[ \\grad_I, \\grad_J ] V_K = R_{IJKL} V^L$. Finally, it is useful to define covariant versions of the derivatives with respect to coordinate time $t$, conformal time $\\eta = \\int \\d t / a$ and e-folds $N = \\int H \\, \\d t$ as \\begin{equation} \\Dt = \\frac{\\d \\phi^I}{\\d t} \\grad_I , \\qquad \\Deta = \\frac{\\d \\phi^I}{\\d \\eta} \\grad_I , \\qquad \\DN = \\frac{\\d \\phi^I}{\\d N} \\grad_I . \\end{equation} ",
        "conclusions": "\\label{sec:conclusion}  In this paper we have computed the covariant three-point function near horizon-crossing for a collection of slowly-rolling scalar fields with a nontrivial field-space metric. After making a conformal transformation, this framework is sufficiently general to include scenarios where one or more fields are nonmimimally coupled to the Ricci scalar. The subsequent superhorizon evolution can be expressed using a version of the `separate universe' approach.  We concentrate on the broad class of models described by a $\\sigma$-model Lagrangian $\\mathcal{L} = X + V$, where $2X = - \\G_{IJ} \\partial_\\mu \\vp^I \\partial^\\mu \\vp^J$, and obtain expressions for the two- and three-point functions. The presence of a nontrivial field-space metric leads to technical subtleties. First, to obtain a covariant formalism we must be careful to define perturbations as tangent-space vectors $Q^I$ using the method of Gong~\\&~Tanaka. When computing correlation functions involving the $Q^I$, the Feynman diagram expansion introduces explicit factors of the `trajectory propagator' which implements parallel transport along the inflationary trajectory.  Second, new interaction vertices appear which involve explicit factors of the Riemann curvature tensor. Therefore the two- and three-point functions receive modifications of two types. The first follow from promotion of the flat field-space perturbation $\\delta \\vp^I$ to $Q^I$ and covariantize the result for $\\G^{IJ} = \\delta^{IJ}$. As in general relativity, covariantization is achieved by exchanging partial derivatives for covariant derivatives and contracting all indices with the field-space metric. The second involve explicit factors of the field-space Riemann curvature. These modify both quantum interference effects operating near horizon exit, and the interactions between growing modes far outside the horizon.  In~\\S\\ref{sec:evolution} we developed a covariant version of the `separate universe' formalism to account for these superhorizon-scale interactions. The covariant Jacobi equation~\\eqref{eq:jacobi-eq} automatically incorporates curvature contributions which influence the evolution of the two- and three-point function. We have shown that this correctly reproduces the time-dependent growing modes near horizon-crossing generated by the apparatus of quantum field theory. In particular, it matches the two lowest-order divergences at single-logarithm order and the leading divergence at double-logarithm order.  The Jacobi approach leads to covariant `time evolution operators' ${\\bigamma^I}_m$ and ${\\bigamma^I}_{(mn)}$ which can be obtained straightforwardly by direct integration of~\\eqref{eq:gamma2-eq}--\\eqref{eq:gamma3-eq}. Together with the covariant gauge transformations derived in~\\S\\ref{sec:observables} these yield covariant `$\\delta N$ coefficients' $\\biN_m$ and $\\biN_{(mn)}$ which define the `separate universe' expansion of the curvature perturbation $\\zeta$ in Eq.~\\eqref{eq:covariant-deltaN}. This provides a clear and economical framework enabling perturbations to be evolved in a slow-roll inflationary model with nontrivial field-space metric.  We always retain the option to abandon manifest covariance and work with the coordinate variation $\\delta\\vp^I$. The traditional separate-universe expansion for $\\delta\\vp^I$ is unchanged by the presence of a nontrivial field-space metric, and our predictions for the autocorrelation functions of $\\zeta$ cannot vary because $\\zeta$ is a field-space scalar. The advantage of the covariant formulation is one of convenience and practicality.  The phenomenology of the new Riemann-tensor terms may be interesting. In a canonical scenario, interactions among superhorizon modes are suppressed by three powers of $\\dot{\\phi}/H$, and are therefore relatively slow. (Note that one should not regard this suppression as an indication of how large the bispectrum can become. Rather, it is an indication of the \\emph{timescale} over which it can evolve.) However, Eq.~\\eqref{eq:bispectrum} shows that curvature-mediated interactions are suppressed by only a single power of $\\dot{\\phi}/H$. In a model where the field-space curvature is $\\Or(1)$ these could lead to much more rapid evolution. It will be interesting to study these effects in more detail, and we hope to return to this in future work."
    },
    "1208/1208.4478_arXiv.txt": {
        "abstract": "{ High-resolution soft X-ray spectroscopic observations of single hot white dwarfs are scarce. With the \\textsl{Chandra} Low-Energy Transmission Grating, we have observed two white dwarfs, one is of spectral type DA (\\object{LB\\,1919}) and the other is a non-DA of spectral type PG\\,1159 (\\object{PG\\,1520+525}). The spectra of both stars are analyzed, together with an archival \\textsl{Chandra} spectrum of another DA white dwarf (\\object{GD\\,246}). }{ The soft X-ray spectra of the two DA white dwarfs are investigated in order to study the effect of gravitational settling and radiative levitation of metals in their photospheres. \\object{LB\\,1919} is of interest because it has a significantly lower metallicity than DAs with otherwise similar atmospheric parameters. \\object{GD\\,246} is the only white dwarf known that shows identifiable individual iron lines in the soft X-ray range. For the PG\\,1159 star, a precise effective temperature determination is performed in order to confine the position of the blue edge of the GW~Vir instability region in the HRD. }{ The \\textsl{Chandra} spectra are analyzed with chemically homogeneous as well as stratified NLTE model atmospheres that assume equilibrium between gravitational settling and radiative acceleration of chemical elements. Archival EUV and UV spectra obtained with \\textsl{EUVE}, \\textsl{FUSE}, and \\textsl{HST} are utilized to support the analysis. }{ No metals could be identified in  LB\\,1919. All observations are compatible with a pure hydrogen atmosphere. This is in stark contrast to the vast majority of hot DA white dwarfs that exhibit light and heavy metals and to the stratified models that predict significant metal abundances in the atmosphere. For GD\\,246 we find that neither stratified nor homogeneous models can fit the \\textsl{Chandra} spectrum. The \\textsl{Chandra} spectrum of PG\\,1520+525 constrains the effective temperature to \\Teff = $150\\,000 \\pm 10\\,000$\\,K. Therefore, this nonpulsating star together with the pulsating prototype of the GW\\,Vir class (PG\\,1159$-$035) defines the location of the blue edge of the GW\\,Vir instability region. The result is in accordance with predictions from nonadiabatic stellar pulsation models. Such models are therefore reliable tools to investigate the interior structure of GW~Vir variables. }{ Our soft X-ray study reveals that the understanding of metal abundances in hot DA white dwarf atmospheres is still incomplete. On the other hand, model atmospheres of hydrogen-deficient PG\\,1159-type stars are reliable and reproduce well the observed spectra from soft X-ray to optical wavelengths. } ",
        "introduction": "}  An unexpectedly small number of white dwarfs (WDs) were detected in the \\textsl{ROSAT} PSPC X-ray all sky survey \\citep{1996A&A...316..147F}. Most of them (161) are of spectral type DA (pure-hydrogen optical spectra), with an additional three DOs (helium-dominated), three DAOs (mixed H/He optical spectra), and eight PG\\,1159 stars (He--C--O dominated). It was realized that the atmospheric opacity of radiatively levitated metals effectively blocks the outward leakage of X-rays in the vast majority of DAs with \\Teff $>$ 40\\,000\\,K \\citep{1997MNRAS.286...58B}. Furthermore, the interstellar medium proved denser than expected in many lines of sight (see, e.g., \\citealt{1999A&A...352..308W}), thus affecting cooler WDs.  High-resolution soft X-ray spectroscopy of hot WDs enables the identification of chemical species in their photospheres that cannot be detected in other wavelength ranges. However, such observations are rather scarce and only became feasible with the advent of the \\textsl{Chandra} observatory with its Low-Energy Transmission Grating (LETG).  Until recently, only five WDs were observed with \\textsl{Chandra} LETG. \\object{H\\,1504+65} is an extremely hot (\\Teff= 200\\,000\\,K), peculiar PG\\,1159 star, obviously a naked C--O or O--Ne WD, whose soft X-ray spectrum is characterized by highly ionized O, Ne, and Mg lines \\citep{2004A&A...421.1169W}. On the other hand, \\object{GD\\,246} is a DA, and its \\textsl{Chandra} spectrum allowed the first unambiguous identification of lines from highly ionized iron \\citep{2002ASPC..262...57V}. The \\textsl{Chandra} spectra of the two DAs \\object{HZ\\,43} and \\object{Sirius~B} are perfectly matched by pure-hydrogen atmospheres and are used as soft-X-ray calibration targets \\citep{2000SPIE.4012..700P,2002ASPC..262...57V,2006A&A...458..541B}. A \\textsl{Chandra} observation of the hot DO \\object{KPD\\,0005+5106} (\\Teff= 200\\,000\\,K) was used to prove that the soft X-ray emission is of photospheric and not coronal origin; however, the S/N ratio is insufficient to identify individual spectral lines \\citep{2005ApJ...625..973D}.  Few other WDs are bright enough in soft X-rays to obtain useful \\textsl{Chandra} spectra. We have observed two of them (\\object{LB\\,1919} and \\object{PG\\,1520+525}) and including archival data of \\object{GD\\,246} present a spectral analysis in this paper. The DA \\object{LB\\,1919} is of interest because previous \\textsl{Extreme Ultraviolet Explorer (EUVE)} spectroscopy showed an unexpectedly high soft X-ray flux that reveals an unexplained low metal abundance \\citep{1998AuAfortable...329.1045W}, in contrast to expectations from radiative-levitation theory. We performed X-ray spectroscopy in order to identify individual elements that could hint at the origin of the metal deficiency.  We observed \\object{PG\\,1520+525} in order to constrain its effective temperature. The motivation lies in the fact that this nonpulsating PG\\,1159 star, together with the pulsating prototype \\object{PG\\,1159--035}, defines the blue edge of the GW\\,Vir instability region in the Hertzsprung-Russell diagram (HRD), see e.g. \\cite{2007A&A...462..281J}. The homogeneous atmosphere of \\object{PG\\,1520+525} provides a convenient comparison for the investigated DA white dwarfs because it is not subject to the uncertainties in the treatment of radiative levitation and stratification of atmospheric composition. Thus, \\object{PG\\,1520+525} provides a test case for the model atmospheres employed in this study.  A failure to achieve a consistent fit to the UV/optical and X-ray spectra of the homogeneous PG1159 photosphere would indicate serious model deficiencies in addition to problems that would be encountered with radiative levitation physics in the case of hot DA WDs.  This paper is organized as follows. We first describe in more detail the motivation of our soft X-ray analyses (Sect.\\,\\ref{sect:softx}). Then we specify our model atmosphere calculations and the atomic data used (Sect.\\,\\ref{sect:models}). In Sect.\\,\\ref{sect:obs}, we report on the observations utilized in our analysis. In Sections \\ref{sect:lb} to \\ref{sect:pg} we delineate the analysis procedure of our three program stars one at a time, and in Sect.\\,\\ref{sect:sum} we conclude. ",
        "conclusions": "\\label{sect:sum}  We analyzed the \\textsl{Chandra} spectra of three hot WDs. Two of them are H-rich WDs and the other one is a H-deficient PG\\,1159 star.  \\subsection{The DA white dwarfs \\object{LB\\,1919} and \\object{GD\\,246}}  The primary aim of our \\textsl{Chandra} observation of LB\\,1919 was to explain its relatively low metallicity compared to similar objects as suggested by previous analyses of \\textsl{EUVE} spectra. It turned out that no metal features were detected in the \\textsl{Chandra} spectrum. The same result was found from our analysis of a \\textsl{FUSE} spectrum. In essence, all data are compatible with the assumption that LB\\,1919 has a pure hydrogen atmosphere. This is in conflict with our stratified models. In them, the vertical run of individual element abundances is computed from the assumption of equilibrium between gravitational downward pull and radiative upward acceleration, i.e., the metal abundances are not free parameters but computed as functions of \\Teff\\ and \\logg. They predict that significant amounts of light and heavy metals should be accumulated and readily detectable in the atmosphere of LB\\,1919.  A few other hot DAs with similar parameters that also do not have detected metals are known \\citep{2003MNRAS.341..870B}; one famous example is HZ\\,43A (\\Teff = 51\\,000\\,K, \\logg = 7.9, \\citealt{2006A&A...458..541B}). The reason for the purity of their atmospheres is completely unknown. It may be speculated that these peculiar DAs have no heavy-element reservoir that is assumed to be present in the radiative levitation calculations. \\citet{2003MNRAS.341..870B} consider whether the possible depletion of those reservoirs by selective mass-loss could be responsible for that phenomenon or whether the progenitors were metal-poor. They argue that, for radiation-driven winds, mass loss should be lower in the pure H stars than in those containing heavy elements, because the wind is driven by metal lines. Furthermore, these WDs are local disc objects, so it seems unlikely that any of the progenitors could have been metal poor.  We have also analyzed an archival \\textsl{Chandra} observation of the hot DA GD\\,246. Our work was motivated by the fact that this is the only DA that exhibits individual metal-line features in the soft X-ray range. As for LB\\,1919, we utilized chemically homogeneous as well as stratified models. While we in principle expect that the stratified models are a better representation, the \\textsl{Chandra} observation cannot be fitted satisfactorily by either type of model. A better fit by stratified models is only achieved when nickel is removed artificially. Otherwise, the atmospheric opacity becomes much too large because of a strong overprediction of nickel. In comparison with the \\textsl{FUSE} spectrum, most metals are overpredicted, but at the same time, another one (Ge) is underpredicted. In the latter case, one can speculate that the model atoms are still too small, not having enough line transitions to absorb photon momentum. The overprediction of most metals, however, is more of a problem. We can only speculate that an additional physical mechanism, ignored by our models, is affecting the equilibrium abundances, e.g., (selective) mass-loss or weak magnetic fields.  The complete failure of our model atmospheres, homogeneous or stratified, to fit the \\textsl{EUVE} observation reinforces the conclusion that some physics is missing. The same conclusion was drawn by \\citet{2003MNRAS.341..870B} from the fact that UV spectra from DAs with similar \\Teff\\ and \\logg\\ display a wide variety of metal abundances. These authors considered the possibility that accretion of interstellar or circumstellar matter could cause this variety. Recent investigations have revealed the frequent incidence of dust debris disks around WDs (e.g., \\citealt{2011AIPC.1331..193F}) that are composed of disrupted planetary material.  In order to test the influence of accretion of such material on the observed photospheric abundances, this effect would have to be included in the diffusion/levitation models in a manner performed (e.g., by \\citealt{2009A&A...498..517K}) in cooler WDs in which radiative levitation is negligible. However, it is difficult to understand how an external supply of material would resolve the problem of present diffusion models generally overpredicting metal abundances.  Our analysis of the two DAs comprised the derivation of effective temperature and gravity from the Lyman lines, and the resulting parameters were compared to published results from Balmer-line analyses. We confirm previous findings that conflicting results are obtained from the UV and optical line profile analyses. In particular, Balmer-line temperatures are often significantly lower than Lyman-line temperatures \\citep{2003MNRAS.344..562B}. We found that LB\\,1919 is another rare case where the opposite effect was found. While this obviously points to a shortcoming in the WD atmosphere models, there is no indication as to which physical ingredient is treated inadequately or is missing. These uncertainties in the derived atmospheric parameters add to the problems in a detailed quantitative comparison, element by element, of observed metal abundances with predictions from radiative levitation models.  \\subsection{The PG\\,1159 star \\object{PG\\,1520+525}}  The \\textsl{Chandra} spectrum of the nonpulsator PG\\,1520+525 was used to constrain its effective temperature and to compare it with that of the pulsator PG\\,1159$-$035. They both confine empirically the blue edge of the GW~Vir instability region for a particular envelope composition. The position of the edge predicted by the nonadiabatic pulsation models of \\citet{0067-0049-171-1-219} is consistent with the spectroscopic results. This is strong proof for the predictive power of these  models. The interior structure of PG\\,1159 stars that is inferred by their usage appears to be very reliable, so that corresponding asteroseismologic analyses are based on solid ground.  \\begin{figure} \\centering \\includegraphics[scale=0.45]{19718fig24.ps} \\caption{Pulsating (filled circles) and nonpulsating (empty circles) PG\\,1159 stars in the $\\log \\Teff - \\logg$ diagram. Evolutionary tracks are labeled with the respective stellar masses in M$_{\\sun}$ \\citep{2006A&A...454..845M}. The red edge (short dashed line) of the instability region \\citep{2006MmSAI..77...53Q} and two blue edges (long dashed lines,  \\citealt{0067-0049-171-1-219}) are shown. The upper and lower blue edges are for $z=0$ and $0.007$, respectively.} \\label{pulsators_inst_neu} \\end{figure}"
    },
    "1208/1208.4364_arXiv.txt": {
        "abstract": "We reconstruct $\\gamma$-ray opacity of the universe out to $z\\lsim 3-4$ using an extensive library of \\NLF\\ observed galaxy luminosity function (LF) surveys extending to high redshifts. We cover the whole range from UV to mid-IR (0.15-25\\mic) providing for the first time a robust empirical calculation of the $\\gamma\\gamma$ optical depth out to several TeV. Here, we use the same database as \\citet{Helgason12} where the extragalactic background light was reconstructed from LFs out to 4.5\\mic\\ and was shown to recover observed galaxy counts to high accuracy. We extend our earlier library of LFs to 25\\mic\\ such that it covers the energy range of pair production with $\\gamma$-rays (1) in the entire {\\it Fermi}/LAT energy range, and (2) at higher TeV energies probed by ground-based Cherenkov telescopes. In the absence of significant contributions to the cosmic diffuse background from unknown populations, such as the putative Population III era sources, the universe appears to be largely transparent to $\\gamma$-rays at all {\\it Fermi}/LAT energies out to $z\\sim 2$ whereas becoming opaque to TeV photons already at $z\\lsim 0.2$ and reaching $\\tau \\sim 10$ at $z=1$. Comparison of the currently available {\\it Fermi}/LAT gamma-ray burst and blazar data shows that there is room for significant emissions originating in the first stars era. ",
        "introduction": "The extragalactic background light (EBL) supplies opacity for propagating high energy GeV-TeV photons via an electron-positron pair production ($\\gamma\\gamma\\rightarrow e^+e^-$) \\citep{Nikishov62,Stecker71}. Determining the transparency of the universe is of fundamental importance for a wide variety of current observatories such as the space-borne {\\it Fermi}/LAT instrument operating at energies $\\lsim 250-300$ GeV to ground-based $\\gamma$-ray telescopes probing energies $\\gsim$1 TeV. The distance at which the optical depth due to this interaction is $\\tau\\sim 1$ defines a horizon of the observable universe at $\\gamma$-ray energies, and has been a subject of extensive efforts designed to model the build-up of EBL with time from the posited emission history of galaxy populations \\citep[e.g.,][]{Stecker06,Franceschini08,Kneiske&Dole10,Dominguez11}.  In this Letter we show that, with observed galaxy populations over a wide range of wavelengths, one can uniquely reconstruct the optical depth of the universe at these energies out to redshifts $z \\sim 4$. This empirical reconstruction relies exclusively on {\\it data} from an extensive library of galaxy luminosity functions (LFs) encompassing 18 finely sampled wavelengths from UV to mid-IR (0.15-24 \\mic) relevant for the pair-production opacity. This methodology enables robust calculation of the $\\gamma$-ray opacity in the {\\it Fermi}/LAT energy range using galaxy surveys probing $\\lambda \\leq 4.5 $\\mic\\ out to $z\\lsim 4$. Extending to TeV energies, probed by the ground-based Cherenkov observations, we use measurements out to 24 \\mic; this extrapolation is robust for the redshifts currently probed these observations.  This heuristic reconstruction using the {\\it observed} galaxy populations defines the absolute floor of the photon-photon optical depth due to known galaxy populations and deviations from it would allow the characterization of any emissions inaccessible to direct telescopic studies \\citep{Kashlinsky05a,Gilmore12b}.  We use the methodology developed in \\citet{Helgason12} of reconstructing the EBL from observed galaxy populations in a compilation of \\NLF\\ measured LFs covering the UV, optical and near-IR; that compilation is slightly updated compared to Table 1 of \\citet{Helgason12}. The wealth of galaxy survey data has recently reached adequate redshift coverage to make such empirical estimation of the evolving EBL feasible and the reconstruction was shown to reproduce independent data from galaxy counts and the cosmic infrared background \\citep[CIB, ][]{Kashreview}. Our approach is completely independent of theoretical modeling describing the evolution of galaxy populations in that we use the LF data directly observed at all wavelengths out to $z\\sim3-8$ in this heuristic reconstruction from which we derive the optical depth due to pair production \\citep[see also][]{Stecker12}.  Standard cosmological parameters are used below: $\\Omega_{\\rm tot}=1$, $\\Omega_{\\rm m}=0.3$, $H_0=70 {\\rm km \\cdot s^{-1} Mpc^{-1}}$. ",
        "conclusions": "We have shown that it is possible to robustly reconstruct the evolving EBL in the universe using our earlier library of multiwavelength survey data now updated to extend from the UV out to the mid-IR \\citep{Helgason12}. This reconstruction uniquely defines the $\\gamma$-ray opacity out to TeV energies for sources at $z\\lsim 4$ and shows that at the energy bands probed by {\\it Fermi}/LAT, the universe is fairly transparent out to $z\\sim 2-3$, unless unknown sources at high redshifts contribute non-negligible amounts of CIB. At TeV energies, probed by ground based telescopes, the universe becomes optically thick at $z\\sim 0.5$ so any such photons associated with the sources at higher redshifts would have to be of secondary origin. %  Our reconstructed EBL and optical depths are available upon request. This work was supported by NASA Headquarters under the NASA Earth and Space Sciences Fellowship Program - Grant NNX11AO05H. KH is also grateful to {\\it The Leifur Eiriksson Foundation} for its support. We thank W. McConville, B. Magnelli and M. Ricotti for useful communications."
    },
    "1208/1208.4687_arXiv.txt": {
        "abstract": "{Understanding the growth of the cores of giant planets is a difficult problem. Recently, Lambrechts \\& Johansen (2012; LJ12) proposed a new model in which the cores grow by the accretion of pebble-size objects, as the latter drift towards the star due to gas drag.}{We investigate the dynamics of pebble-size objects in the vicinity of planetary embryos of 1 and 5 Earth masses and the resulting accretion rates.}{We use hydrodynamical simulations, in which the embryo influences the dynamics of the gas and the pebbles suffer gas drag according to the local gas density and velocities.}{The pebble dynamics in the vicinity of the planetary embryo is non-trivial, and that it changes significantly with the pebble size. Nevertheless, the accretion rate of the embryo that we measure is within an order of magnitude of the rate estimated in LJ12 and tends to their value with increasing pebble-size.}{The model by LJ12 has the potential to explain the rapid growth of giant planet cores. The actual accretion rates however, depend on the surface density of pebble size objects in the disk, which is unknown to date.} ",
        "introduction": "The formation of the massive cores of giant planets within the short timescale allowed by the survival of a proto-planetary disk of gas and solids (a few My; Haisch et al. 2001) is still an open problem. In the classical view, these cores form by collisional coagulation from a disk of km-sized planetesimals, through the well-known processes of {\\it runaway} (Greenberg et al. 1978; Wetherill \\& Stewart 1989) and {\\it oligarchic} growth (Ida \\& Makino 1993; Kokubo \\& Ida 1998). In principle these processes should continue until the largest objects achieve an {\\it isolation mass}, which is a substantial fraction of the initial total mass of local solids. If the initial disk is sufficiently massive (about 10 times the so-called Minimal Mass Solar Nebula or MMSN; Weidenschilling 1977; Hayashi 1981), it is expected that cores of $\\sim 10$ Earth masses ($M_E$) form beyond the so-called {\\it snowline} (The orbital radius beyond which temperature is cold enough that water condenses into ice; Podolak \\& Zucker 2004), as required in the core-accretion model for giant planet formation (Thommes et al. 2003; Goldreich et al. 2004; Chambers 2006).  $N$-body simulations, though, show that reality is not so simple. When the cores achieve a mass of about 1~$M_\\oplus$ they start to scatter the planetesimals away from their neighborhood, instead of accreting them (Ida \\& Makino 1993; Levison et al. 2010), which slows their accretion rate significantly. It has been proposed that gas drag (Wetherill \\& Stewart 1989) or mutual inelastic collisions (Goldreich et al. 2004) prevent the dispersal of the planetesimals by damping their orbital eccentricities, but in this case the cores open gaps in the planetesimal disk (Levison \\& Morbidelli 2007; Levison et al. 2010), like the satellites Pan and Daphis open gaps in Saturn's rings.  Thus the cores isolate themselves from the disk of solids.  This effectively stops their growth. It has been argued that planet migration (Alibert et al. 2004) or the radial drift of sub-km planetesimals due to gas drag (Rafikov 2004) break the isolation of the cores from the disk of solids but, again, $N$-body simulations show that the relative drift of planetesimals and cores simply collects the former in resonances with the latter (Levison et al. 2010); this prevents the planetesimals from being accreted by the cores.  In fact, only planetesimals smaller than a few tens of meters drift in the disk fast enough to avoid trapping in any resonance with a growing core (Weidenschilling \\& David 1985).  In a recent paper, Lambrechts \\& Johansen (2012), hereafter LJ12, have proposed a new model of core growth, which argues that, if the mass in the disk is predominantly carried by pebbles of a few decimeters in size, the largest planetesimals accrete pebbles very efficiently, rapidly growing to several Earth masses (see also Johansen \\& Lacerda 2010; Ormel \\& Klahr 2010 and Murray-Clay et al. 2011).  More specifically, this model builds on the recent planetesimal formation model (Youdin \\& Goodman 2005; Johansen et al. 2006, 2007, 2009) in which large planetesimals (with sizes from $\\sim 100$ up to $\\sim$1,000km) form by the collapse of a self-gravitating clump of pebbles, concentrated to high densities by disk turbulence and the streaming instability.  The mechanism by which, once formed, planetesimals can keep accreting background pebbles is described hereafter.  Pebbles are strongly coupled with the gas; thus they encounter the already-formed planetesimals with a velocity $\\Delta v$ that is equal to the difference between the Keplerian velocity and the orbital velocity of the gas (slightly sub-Keplerian due to the outward pressure gradient). LJ12 define the planetesimal {\\it Bondi radius} as the distance at which the planetesimal exerts a deflection of one radian on a particle approaching with a velocity $\\Delta v$: \\begin{equation} R_B={{GM}\\over{\\Delta v^2}} \\label{Bondi} \\end{equation} where $G$ is the gravitational constant and $M$ is the planetesimal mass (obviously the deflection is larger if the particle passes closer than $R_B$). LJ12 showed that all pebbles with a stopping time $t_s$ smaller than the Bondi time $t_B=R_B/\\Delta v$ that pass within a distance $R=(t_s/t_B)^{1/2} R_B$ spiral down towards the planetesimal and are accreted by it. Thus, the growth rate of the planetesimal is: \\begin{equation} \\dot{M}=\\pi\\rho R^2\\Delta v \\label{Bondi-accrete} \\end{equation} where $\\rho$ is the volume density of the pebbles in the disk.  From (\\ref{Bondi}), the Bondi radius grows with the planetesimal mass. LJ12 also showed that, when the Bondi radius exceeds the scale height of the pebble layer, the accretion rate becomes \\begin{equation} \\dot{M}=2R\\Sigma \\Delta v \\end{equation} where $\\Sigma$ is the surface density of the pebbles. Moreover, when the Bondi radius exceeds the Hill radius $R_H$, the accretion rate becomes \\begin{equation} \\dot{M}=2R_H \\Sigma v_H \\label{Hill-accretion} \\end{equation} where $v_H$ is the Hill velocity (i.e. the difference in Keplerian velocities between two circular orbits separated by $R_H$).  With these formulae, and assuming that $\\Sigma$ stays constant and is close to the nominal density of solids in the MMSN, LJ12 showed that the formation of 10 $M_E$ cores is possible within 1 My essentially anywhere in the disk (up to $\\sim 50 AU$).  There are two main advantages in the LJ12 model. First, it can form 10 $M_E$ cores of giant planets within the lifetime of the disk, a result very difficult to achieve by other models. Second, because the accretion rate (\\ref{Bondi-accrete}) is very sensitive on the planetesimal mass ($\\dot{M}\\propto M^2$), in practice only the largest planetesimals formed in the turbulent model can effectively grow in mass by this process: the minimal mass for triggering significant Bondi accretion (see eq. \\ref{Bondi-accrete}) is about the mass of Ceres in the asteroid belt and about the mass of Pluto in the Kuiper belt. Thus this model explains the maximal sizes observed in the asteroid and Kuiper belt populations. In essence, in this model bodies smaller than Ceres (respectively Pluto for the Kuiper belt) remained small bodies (the asteroids and KBOs we see today), whereas those bigger than this threshold kept accreting pebbles and became massive objects (embryos) which then were removed by migrating away and (possibly) participating to the build-up of the giant planets. Both these aspects of the model are very appealing.  However the study conducted in LJ12, both in the analytic and in the numerical parts, assumes that the motion of the gas is not perturbed by the planetesimal. This assumption is good for a Ceres-mass planetesimal, accreting as in (\\ref{Bondi-accrete}), but it is far from reality for planetary embryos (Earth mass or larger), accreting through their Hill sphere as in (\\ref{Hill-accretion}). In fact, these objects modify the gas streamlines significantly: a spiral density wave is formed in the disk and the gas near the orbit of the planet has horseshoe motion. An over-density of gas is also formed inside the planet's Hill sphere. It is not clear a priori what are the effects of these structures on the pebble accretion rate. This is precisely what we investigate in this paper with more realistic hydro-dynamical simulations.  In section 2, we describe our methods: the simulation tool that we have developed and the parameters that we adopt. In section 3 we present our results, for two embryo masses and 4 pebble sizes. Our goal is three-fold: (i) describe and understand the dynamics of the pebbles for the different mass and size cases; (ii) evaluate the accretion rate by the embryo and compare it with the LJ12 estimate and (iii) evaluate the ``filtering factor'', that is the fraction of the pebbles that do not drift by the orbit of the planet because they are accreted by the embryo.  This factor is important. If it is large, of a sequence of embryos radially distributed in the disk, only the outermost one(s) can accrete; instead, if it is small then the full system of embryos can grow, in an oligarchic fashion.  Our conclusions and discussion of a coherent scenario of giant planet formation conclude the paper in section 4. ",
        "conclusions": "In this paper we have tested the scenario of giant planet core formation proposed by LJ12 with hydrodynamical simulations that fully account for (i) the interaction between the growing core and the gas of the disk and (ii) the local drag exerted by the gas on pebbles and boulders.  We have found that the pebble dynamics in the vicinity of the planetary embryo is non-trivial, and that it changes significantly with the pebble size. Nevertheless, the accretion rate of the embryo that we measure is within an order of magnitude of the rate estimated in LJ12 and tends to their value with increasing pebble-size.  The accretion of pebbles can continue until the embryo's mass is of the order of 50 Earth masses (in solids and gas). This may have important implications on the onset of runaway accretion of gas by the growing giant planets and can help explaining the enrichment in solids observed in Jupiter's atmosphere.  The actual accretion rate of an embryo depends on the amount of mass $\\Sigma$ available in pebbles, which is not known a priori, given that pebbles are consumed by the formation of planetesimals and by the accretion of the embryos themselves. Nevertheless, the accretion rates that we find in Fig.~\\ref{accretion} are potentially large. For instance, if a MMSN of solids (20g/cm$^2$ at 1 AU) were available in 20cm pebbles, the mass doubling time for a 1~$M_E$ embryo would be only 5,500 years! This illustrates the importance of the LJ12 model for the growth of giant planets cores.  Given that the LJ12 model is built in the same framework as the model that explains the rapid formation of planetesimals (Johansen et al. 2007), we believe that the community has now, for the first time, a coherent scenario to explain the the early phases of planet growth. However, we do not see any evident reason for which only a small number (4-6) of giant planet cores should grow in the disk, as suggested by the number of giant planets of our solar system, including rogue ice-giants potentially lost during the dynamical evolution that followed planet formation (Nesvorny 2011; Nesvorny \\& Morbidelli 2012). Thus, we think that, most likely, the LJ12 model explains the formation of massive planetary embryos of a few Earth masses, but an additional stage is needed to form the giant planet cores (10 Earth masses or more).  This is the scenario that we envision. The embryos formed by the LJ12 mechanism, once massive enough, start to migrate in the disk due to planet-disk interactions. Recent results on migration in radiatively cooling disks (Paardekooper \\& Mallema 2006; Baruteau \\& Masset 2008; Kley \\& Crida 2008; Paardekooper et al. 2010; Masset \\& Casoli 2010; Bitsch \\& Kley 2011) show that the embryos migrate from all directions toward an orbital radius where migration is canceled out by the compensation of competing torques. This convergent migration towards the same region can promote the mutual accretion of embryos, eventually reducing a system of a large number of embryos into a system of a smaller number of larger objects, i.e. a handful of giant planet cores (Horn et al. 2012). Admittedly, this scenario is still speculative and more work is needed to prove its validity.  We stress, however, that a final phase of core formation characterized by mutual collisions of embryos would explain, in a natural way, the massive impacts that are needed to explain the current obliquities of Uranus and Neptune (Morbidelli et al. 2012)."
    },
    "1208/1208.4585.txt": {
        "abstract": "We present mid-infrared spectra and photometry of thirteen redshift $0.4<z<1$ dust-reddened quasars obtained with {\\it Spitzer} IRS and MIPS. We compare properties derived from their infrared spectral energy distributions (intrinsic AGN luminosity and far-infrared luminosity from star formation) to the host luminosities and morphologies from HST imaging, and black hole masses estimated from optical and/or near-infrared spectroscopy. Our results are broadly consistent with models in which most dust reddened quasars are an intermediate phase between a merger-driven starburst triggering a completely obscured AGN, and a normal, unreddened quasar. We find that many of our objects have high accretion rates, close to the Eddington limit. These objects tend to fall below the black hole mass -- bulge luminosity relation as defined by local galaxies, whereas most of our low accretion rate objects are slightly above the local relation, as typical for normal quasars at these redshifts. Our observations are therefore most readily interpreted in a scenario in which galaxy stellar mass growth occurs first by about a factor of three in each merger/starburst event, followed sometime later by black hole growth by a similar amount. We do not, however, see any direct evidence for quasar feedback affecting star formation in our objects, for example in the form of a relationship between accretion rate and star formation. Five of our objects, however, do show evidence for outflows in the [\\ion{O}{3}]5007\\AA~ emission line profile, suggesting that the quasar activity is driving thermal winds in at least some members of our sample. ",
        "introduction": "The study of quasars across cosmic time has proven to be important as their evolution and growth is intimately linked to that of their host galaxy. Mergers have long been invoked as the primary ignition mechanism of quasar activity, at least for the high luminosity objects \\citep{sanders96,redqso-hst,nicola}. The loss of angular momentum during a gas rich merger allows the gas to be funneled to the center igniting a starburst and fueling the central black hole. In this model, eventually the quasar grows strong enough and develops a feedback mechanism that expels the obscuring material and shuts off star formation \\citep{silkrees,hopkins08}. This scenario is further supported by the existence of tight correlations between the properties of galaxy bulges and their central black holes \\citep{magorrian,ferrarese}.  Putting the picture together observationally is challenging, however. Populations of quasars selected via different techniques, although overlapping, differ substantially in their detailed properties and may bias towards a particular type of quasar or evolutionary phase e.g. AGN selected in the infrared may belong to different class or evolutionary phase than AGN selected in the radio or optical regime \\citep{hickox09}. In an ideal survey we would be able to select quasars at specific points in their lifetimes. For example, to test a starburst-AGN connection we would prefer to select young, recently ignited objects to see if there is still star formation occurring in the host galaxies, and to identify likely feedback mechanisms.  One approach to finding candidates for young quasars is to search for dust obscured quasars that may correspond to objects yet to fully clear out the dust and gas surrounding them. Conventionally, obscured quasars are divided into two classes: (1) the so-called type-2 quasars - objects showing only narrow emission lines with large inferred extinctions ($A_V > 5-100$) towards the nucleus, most likely caused by a torus surrounding the accretion disk as required in the AGN  ``unification by orientation'' model \\citep[e.g.,][]{antonucci,urry} and (2) moderately reddened quasars ($A_V \\sim 1-5$), which still show broad emission lines in the rest-frame optical, and whose continuum is still dominated by the quasar rather than the host galaxy. These latter objects are obscured by a cold absorber along the line of sight to the quasar, most likely in the host galaxy. They may thus represent the young objects in the final stages of emerging from their dusty cocoons. Throughout this paper the terms red, dust-obscured and moderately reddened quasars all refer to the latter category of objects.  Dust obscured quasars of all types are missing from, or are severely underrepresented in optical surveys, but are present in the radio \\citep[e.g.,][]{postman}, in the X-ray \\citep[e.g.,][]{hickox07}, when selected on the basis of narrow optical AGN emission lines \\citep{zakamska03}, or when selected in the near-infrared \\citep{cutri}. Mid-infrared surveys \\citep[e.g.,][]{lacy04,lacy07a,stern05,alonso06,donley07} have also been extremely successful at finding and identifying obscured quasars over a wide range of reddenings, redshifts, and luminosities. Surveys combining deep radio and mid-infrared data show that the optically- and/or X-ray-selected quasar population constitutes less than half of the total population of quasars \\citep{donley05,martinez05}. Joint selection using radio and near-infrared, as used for the sample in this paper, has proven to be one of the most reliable ways to select the moderately obscured quasar population. The requirement of a bright radio source and very red optical through near-infrared colors results in a set of candidates with relatively little contamination from normal galaxies, low luminosity AGN and stars \\citep{gregg02,lacy02,eilat04,f2m07,glikman12,f2ms}.  The success of the {\\it Spitzer Space Telescope} \\citep{spitzer} has allowed for the detailed study of quasars in the mid-infrared. Programs with the Infrared Spectrograph \\citep[IRS;][]{irs}, in particular, have refined our knowledge of QSO spectral energy distributions (SEDs) in the mid-infrared. For example, the detection of silicate in emission in IRS spectra of quasars is a strong support for the unification model \\citep{hao05,siebenmorgen05}. Average properties of classes of objects using IRS spectra \\citep[e.g,][]{spoon07} have allowed us to constrain the physical properties of dust in AGN \\citep[e.g.,][]{nikutta09}. On average the spectra of luminous quasars are flat, show little or no PAH emission and the silicate features are in emission, associated with dust re-emission. In contrast, the spectra of Ultraluminous Infrared Galaxies (ULIRGs) show a steep rise towards the long wavelengths, moderate PAH emission and silicate absorption troughs associated with embedded star formation \\citep{hao07}. However, some type 2 and some reddened quasars at moderate to high redshifts show deep silicate absorption features \\citep{lacy07b,martinez08,zakamska08} in contrast to samples of X-ray selected Type-2 quasars, in which they are generally weak \\citep{sturm06}.  The Multiband Imaging Photometer for SIRTF \\citep[MIPS;][]{mips} on board {\\it Spitzer} has also allowed us to sample the colder part of the galaxy SED. While AGN dominate the bright 24$\\mu$m population \\citep{donley08}, the {\\it Spitzer} 70$\\mu$m field population is dominated by ULIRGs which have a high merger fraction \\citep{kar10a,kar10b}. However, AGN and quasars still make up a significant fraction of 70$\\mu$m sources, and high star formation rates in quasar hosts are common \\citep[e.g.][]{quest1,quest2}. Consistent with those results \\cite{shi09} find that Type-1 quasars show star formation rates higher than the field galaxies, with luminosities typical of Luminous Infrared Galaxies (LIRGs, $L_{IR} = 10^{11-12} L_{\\odot}$). In particular, dust reddened  quasars have been found to have higher than usual 60/12$\\mu$m luminosity ratios, an indicator of higher star formation rates \\citep{georgakakis}.  In this paper we will describe {\\em Spitzer} IRS and MIPS observations of 13 redshift $0.4<z<1$ reddened quasars selected using joint radio and near-infrared selection for which we have also obtained Hubble Space Telescope (HST) observations \\citep{redqso-hst}. We use these observations to estimate the intrinsic luminosities of the quasars, and to estimate or constrain the star formation rates in the host galaxies. We then evaluate the evidence that this population corresponds to an intermediate stage in quasar evolution by estimating black hole masses and accretion rates, comparing our HST host galaxy luminosities and morphologies to the other properties, and by searching for evidence of feedback processes at work in the hosts. Throughout this paper we adopt a flat Universe, $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.7$ cosmology.  \\begin{deluxetable*}{lccccc} \\tabletypesize\\scriptsize \\tablecaption{{\\it Spitzer} IRS and MIPS observations \\label{observe}} \\tablewidth{0pt} \\tablehead{ \\colhead{Source} & \\colhead{IRS AORID} & \\colhead{MIPS AORID} & \\colhead{$S_{24}$ (mJy)} & \\colhead{$S_{70}$ (mJy)} & \\colhead{$S_{160}$ (mJy)}} \\startdata F2M0729$+$3336 &22386176 &22386432 &12.9$\\pm$3.0 &$<$15       &$<$30  \\\\ F2M0825$+$4716 &22389760 &22386688 &36.1$\\pm$4.7 &85$\\pm$19   & 46$\\pm$11  \\\\ F2M0830$+$3759 &22390016 &22386944 &26.1$\\pm$3.9 &32$\\pm$12   & $<$30  \\\\ F2M0834$+$3506 &22390272 &22387200 &17.6$\\pm$3.4 &26$\\pm$13   & $<$30  \\\\ F2M0841$+$3604 &22390528 &22390528 &11.7$\\pm$3.0 &53$\\pm$14   & $<$30  \\\\ F2M0915$-$2418 &17540352$^{\\dag}$ &22387712 & 87.3$\\pm$7.0 &207$\\pm$24 &73$\\pm$12 \\\\ F2M1012$+$2825 &22390784 &22387968 & 4.9$\\pm$0.8 &$<$15       & $<$30 \\\\ F2M1113$+$1244 &22391040 &22388224 &59.4$\\pm$5.8 &138$\\pm$25  & 67$\\pm$12 \\\\ F2M1118$-$0033 &22391296 &22388480 &24.4$\\pm$4.0 &80$\\pm$24   & 44$\\pm$10 \\\\ F2M1151$+$5359 &22391552 &22388736 &6.9$\\pm$0.7  &20$\\pm$7    & $<$30 \\\\ F2M1507$+$3129 &22391808 &22388992 &9.2$\\pm$2.5  &22$\\pm$11   & $<$30 \\\\ F2M1532$+$2415 &22392064 &22389248 &28.4$\\pm$4.0 &31$\\pm$13   & $<$30 \\\\ F2M1656$+$3821 &22392320 &22389504 &9.7$\\pm$2.4  &15$\\pm$8    & $<$30 \\\\ \\enddata \\tablecomments{All observations made in cycle 4 (Jul 2007 - May 2008) except for F2M0915-2418, whose IRS spectrum was taken in May 2007 as part of program 30121.} \\end{deluxetable*} ",
        "conclusions": "Our aim in this study was to determine whether the infrared properties of these dust-reddened objects were consistent with the merger-driven co-evolutionary model for the growth of SMBH and their host galaxies. In particular, whether these objects are an intermediate stage between a merger-induced starburst and a ``normal'' quasar with an unreddened, blue continuum, with AGN feedback having stopped further star formation in the host. Our results seem to be consistent with this overall picture, but we see considerable variation in properties such as silicate absorption depth, star formation rate and AGN accretion rate on an object-by-object basis.  \\begin{figure*} \\plotone{f6.eps} \\caption{{\\em Right}: Host galaxy morphological distortion, parameterized as $D$, the ratio of the Gini Coefficient (G) to the Concentration Index (C) against Eddington ratio. Normal galaxies have a mean $D\\approx 1.2$ (Abraham et al.\\ 2003), objects with values below this (shaded in the plot) are relatively undisturbed, those with higher values have increasing amounts of disturbance. Symbol colors as for Figure \\ref{mbhlstar}. {\\em Left:} Morphological distortion versus the log of the ratio of far-infrared luminosity to host luminosity in $B$-band, an indicator for star formation. While there could be a hint of a correlation between the amount of host disturbance and the accretion rate, there is no correlation found for star formation versus the host disturbance.\\label{morphcorr}} \\end{figure*}  Much of the object to object variation is due to the heterogeneity of our sample; some have high reddenings, morphological disturbances, some do not. While they all are luminous quasars, they do span a large range in luminosities, too. In the overall picture, however, trends are evident. If an object is morphologically disturbed, it is much more likely to have a high Eddington ratio and to show Silicate in absorption (be reddened by cold dust). Similarly, the two quasars which did not show any signs of merger interaction in \\cite{redqso-hst} also do not display high Eddington ratios or Silicate in absorption and therefore should not be included characterized as ``young quasars''. Not all dust-obscured quasars fit in the co-evolutionary scenario, but statistically, it is the most likely explanation and therefore our conclusions are to be taken as such.  All these objects have mid-infrared (AGN) luminosities in excess of their far-infrared (presumed starburst) luminosities. This dominance of the AGN suggests that AGN feedback processes, if they are indeed important, will have started to have their impact on star formation in the host galaxy as we would have expected the starburst component to have dominated otherwise. We do not see any obvious quasar feedback ``smoking guns''; broad absorption line features from AGN thermal winds, for example, are not present in any of our objects which are of high enough redshift to see \\ion{Mg}{2} absorption ($z\\stackrel{>}{_{\\sim}}$0.7). However, experience from \\cite{f2ms}, has shown that for the red quasars selected using our criteria the spectrum below $\\approx$ 5000 \\AA~is so extinguished that even deep absorption features are difficult to detect in objects with $z\\stackrel{<}{_{\\sim}}$0.9. Only three of our objects (F2M0729+3336, F2M1012+2825 and F2M1507+3129) have appropriate redshifts, and F2M0729+3336 is too noisy to detect absorption features. Nevertheless, an excess of low-ionization, broad absorption line quasars has been found in our full sample of red quasars, \\cite{f2ms}, and a tentative relationship between wind properties and star formation seen has been seen by \\cite{farrah11} in samples of reddened quasars.  Blueshifted wings to [\\ion{O}{3}]5007\\AA, with velocities relative to the line peak of up to $\\sim$1000 km s$^{-1}$ are seen in five of our thirteen objects (F2M0729+3336, F2M0825+4716, F2M0915+2418, F2M1118-0033, F2M1151+5359). Of the four of these for which we can estimate accretion rates, three have accretion rates just exceeding the Eddington limit, possibly indicating a link with extremely high accretion rates, but the star formation rates in the hosts vary widely. We also lack any information on the spatial extent of these outflows, which may be confined to the nuclear regions. Blueshifted [\\ion{O}{3}] line components are common in compact radio sources: \\cite{htm08} propose that these components are due to outflows of gas driven by the radio jets. The radio sources in our quasars are however significantly less luminous than those in the radio galaxies studied by \\cite{htm08}. Radio-quiet type-2 luminous quasars at low redshift also tend to show outflows in [\\ion{O}{3}]5007\\AA~emission \\citep{greene11a}, suggesting that these are not confined to radio sources, though the most spectacular object in the \\cite{greene11a} sample, J1356+1026, is a relatively strong radio source \\citep{greene11b}. Clearly the role of radio jets in powering these outflows needs to be established.  In agreement with the evolutionary model, we see a large fraction of objects with high black hole accretion rates, consistent with the idea that quasar accretion rates are high in the early phase of growth. The objects with the highest accretion rates also tend to fall below the local black hole mass - bulge luminosity relation by $\\sim 0.5$dex; none of our objects with lower accretion rates are found there. This offset from the local relation is consistent with a picture in which the black hole is yet to grow to its equilibrium size following a major merger. It is also broadly consistent with numerical simulations of black hole growth in quasar merger scenarios, in which the total black mass hole grows by about 0.5 dex with every major merger event \\citep[e.g.,][figure 10]{li07}. \\cite{sarria10} also find that their luminous, dust-obscured quasars at $z\\sim 2$ fall below the black hole mass -- bulge mass relation at that redshift. An interesting exception is F2M0915+2418, which has a high accretion rate, but falls close to the black hole mass -- bulge mass relation. This object also has one of the highest star formation rates in the sample, however, allowing it to evolve along a trajectory that will keep it close to the black hole mass -- bulge mass relation.  We do, however, have some remaining puzzles, where our data do not fit in with the most naive expectations of the evolutionary model. In \\cite{redqso-hst} we saw a relationship between $D$ and the amount of reddening in the host, but in Figure \\ref{morphcorr} there seems to be little correlation between the degree of disturbance of the host galaxy and the star formation rate, except that all high accreting quasars are also disturbed to some degree. There may, however, be a very weak trend for objects with high accretion rates relative to the Eddington Limit to have high values of $D$, see Figure \\ref{morphcorr}. Similarly, there is no clear morphological trend related to the offset of the quasar from the local galaxy luminosity -- black hole mass relation.  To some extent, the lack of clear trends here may be due to an admixture of objects reddened by foreground galaxies. In particular the low accretion rate objects may include some objects where the reddening is from an intervening galaxy through a chance alignment, perhaps from a galaxy in the same group or cluster. F2M1532+2415 is a good example of this, with an apparently highly-disturbed, early stage merger morphology, but both a low accretion rate and a low star formation rate, and F2M0834+3506 is a very good candidate for a low accretion rate broad-line radio galaxy reddened by a foreground irregular galaxy.  In summary, our observations are consistent with the simplest evolutionary models, where a merger induced starburst triggers a obscured quasar which begins life accreting at close to the Eddington rate and later evolves into an unobscured quasar accreting at the $\\sim 0.1$ of the Eddington rate more typically observed in the normal quasar population. We see no obvious evidence, however, of feedback affecting star formation in this particular sample (though we do see evidence of strong outflows of ionized gas in some objects). Nor do we see a clear sequence or progression from a newborn, highly obscured, high accretion rate AGN with a very disturbed host into a quiescent, almost unreddened object with a lower accretion rate. Our quasars were, however, selected within a fairly narrow range of luminosity (due to the 2MASS detection requirement and restricted redshift range) and reddening (if the reddening were too high, they would be classed as type-2s, but if it were too low they would not be picked out as red quasars), so we might not expect to be able to see much of an evolutionary sequence.  These observations are most consistent with a picture in which the majority of objects undergo the starburst/merger phase well before we see these reddened quasars emerge, though not so long after that evidence for morphological disturbance of the hosts has been erased. The evolution in the black hole mass -- bulge mass plane is then a shift along the bulge mass axis by $\\sim +0.5$dex followed $\\sim 10^{8}$yr later by a shift along the black hole mass axis by a similar amount. In one case, however, (F2M0915+2418) we do seem to see the starburst and accretion happening simultaneously. As our selection techniques for the dust-reddened quasar population expand, particularly as mid-infrared selected samples become large enough to include significant numbers of reddened type-1 quasars we expect to be able to expand our samples considerably, and trace out more of the evolutionary scenario."
    },
    "1208/1208.3258_arXiv.txt": {
        "abstract": "The mechanical properties of a neutron star crust, such as breaking strain and shear modulus, have implications for the detection of gravitational waves from a neutron star as well as bursts from Soft Gamma-ray Repeaters (SGRs). These properties are calculated here for three different crustal compositions for a non-accreting neutron star that results from three different cooling histories, as well as for a pure iron crust. A simple shear is simulated using molecular dynamics to the crustal compositions by deforming the simulation box. The breaking strain and shear modulus are found to be similar in the four cases, with a breaking strain of ${\\sim 0.1}$ and a shear modulus of ${\\sim 10^{30}\\;\\mathrm{dyne\\,cm^{-2}}}$ at a density of $\\rho = 10^{14}\\,\\mathrm{g\\; cm^{-3}}$ for simulations with an initially perfect BCC lattice. With these crustal properties and the observed properties of {PSR~J2124-3358} the predicted strain amplitude of gravitational waves for a maximally deformed crust is found to be greater than the observational upper limits from LIGO. This suggests that the neutron star crust in this case may not be maximally deformed or it may not have a perfect BCC lattice structure. The implications of the calculated crustal properties of bursts from SGRs are also explored. The mechanical properties found for a perfect BCC lattice structure find that crustal events alone can not be ruled out for triggering the energy in SGR bursts. ",
        "introduction": "The strength of the material that comprises neutron-star crusts determines the maximal strength of gravitational radiation from a single deformed neutron star \\citep{chamel08}. Furthermore, yielding events of the neutron star crust have also been associated with bursts from Soft Gamma-ray Repeaters (SGRs) \\citep{chamel08}. The neutron star crust consists of the outer regions of the star ranging from a density of $10^6\\mathrm{g \\, cm^{-3}}$ up to nuclear density of $2\\times10^{14}\\mathrm{g\\, cm^{-3}}$ \\citep{lattimer04}. The less dense regions of the crust consist of a lattice of nuclei, with the ground state of a body centred cubic (BCC) lattice \\citep{chamel08}. At higher densities the crust is composed of a mixture of nuclei, electrons, protons and free neutrons. The mechanical properties of the neutron star crust which are important for the consideration of crustal deformation include the breaking strain, or maximal degree of deformation before yielding, the amount of stress associated with this strain, as well as the shear modulus of crustal material.  Initial estimates of the breaking strain of the neutron star crust were made by comparing to terrestrial matter. Arguing that the impurities expected in the neutron star crust would result in weakening the structure, the breaking strain of the crust could range between $\\phi_m = 10^{-5}$ and $\\phi_m = 10^{-2}$ \\citep{smoluchowski70}. The shear modulus of the crust has been predicted by treating the neutron star crust as a Coulomb lattice, which has the approximate relationship of $\\mu \\propto (Ze)^2n^{4/3}$, where the charge $Z$ ranges between 30 and 50, and $n$ is the number density, which leads to a expected shear modulus of $\\mu = 10^{30} \\mathrm{dyne\\, cm}^{-2}$ \\citep{smoluchowski70}.  Recently, with molecular-dynamics simulations, the breaking strain of a system representing an accreted neutron star crust has been calculated. The crustal composition of the accreted crust consists of isotopes ranging from $Z=8$ to $Z = 47$ \\citep{gupta07}. The breaking strain for the accreted crust was found to occur around $\\phi_m = 0.1$, which is larger than had been predicted, and this result was also found to only moderately affected by the introduction of impurities, defects and grain boundaries \\citep{horowitz09b}.  In this work molecular-dynamics simulations are used in order to investigate the crustal properties of a non-accreting neutron star. The molecular-dynamics calculations are carried out with the open source software LAMMPS (Large Atomic/Molecular Massively Parallel Simulator) \\citep{plimpton95}, which is available from Sandia Laboratories\\footnote{http://lammps.sandia.gov}. The breaking strain and shear modulus found with these simulations are used in order to place limits on the strain amplitude for a gravitational wave signal expected for a fully deformed neutron star. Limits are also placed on the fracture size required for the energy in a corresponding SGR burst.  In the following section the implementation and  parameters of the molecular-dynamics simulations are discussed. In Section~\\ref{sec:results} the results of these simulations are presented.  The simulation results are applied to the emission of gravitational waves and SGR bursts in section~\\ref{sec:discussion}. In Section~\\ref{sec:conclusions} we summarize our findings. ",
        "conclusions": "\\label{sec:conclusions} In order to investigate the mechanical properties of a non-accreting neutron star crust we have performed various molecular dynamics simulations  using the software LAMMPS.  With these simulations tests were first performed in order to test the size effects of the simulation box size, as well as  determining the melting temperature of the various crustal compositions. A simple shear was applied to the crustal compositions by deforming the initially perfect BCC lattice simulation box of four different crustal compositions, those appropriate for a neutron star cooled via modified Urca, with a thick crust, and with a thin crust, as well as a pure iron crust. Additional simulations were performed to compare the stress-strain relationship of a system initially with a BCC  lattice structure to that which included defects. The temperature dependence on the the breaking strain as well as  the shear modulus was also investigated with additional simple shear simulations at different temperatures. The possibility of a second yielding event as well as the reversibility of the yielding event were also investigated in the four cases.  The calculation of the breaking strain and shear modulus from the simulations was applied to the context of gravitational wave emission from a neutron star due to a deformation on the neutron star surface, such a mountain.  The mechanical properties for the different crustal compositions, which  include pure and impure compositions, with an initial BCC lattice structure were found to share similar characteristics, this includes the pure and impure simulations. The shear modulus at a density of $\\rho = 10^{14}\\,\\mathrm{g\\,cm^{-3}}$ was found to be around $1.78\\times10^{30}\\,\\mathrm{dyne\\,cm^{-2}}$ and the breaking strain was found to be ${\\sim 0.1}$ in all four cases. This breaking strain measurement is in agreement with that found for an accreted neutron star crust composition \\citep{horowitz09b}. As the shear was only applied in one crystallographic orientation, deforming along a different axis could result in different measured properties.  With these mechanical properties the constraints on the detection of gravitational waves were considered. The properties of three pulsars were considered in order to predict the strain amplitude associated with a neutron star with the calculated mechanical properties. LIGO had placed upper limits on the gravitational wave emission of 78 different pulsars and of the three investigated, one was found to have a strain amplitude larger then the upper limit placed on the neutron star. This indicates that if the crust had these calculated properties and was maximally deformed, then gravitational waves would have already been observed with LIGO from PSR J2124-3358. In the case of bursts from SGRs the fracture lengths required for observed burst energies were found to be within the confines of the neutron star crust.  Crustal material which included defects as compared to that which had an initially perfect BCC lattice structure were found to have differing stress-strain relationships. The defects in the simulations were introduced by first melting the perfect BCC lattice and then cooling the material to a solid. This behaviour of the initially imperfect crystal structure was found to display a similar shear modulus to that of the perfect BCC lattice crystal after the first major yielding event occurs. The effect of defects on an accreted neutron star crust were examined in \\citet{horowitz09b} by introducing six grains to the simulation. The addition of defects in the accreted crust case were found to affect  the stress-strain relationship only moderately. The differences between the defects in the non-accreted and accreted crust composition could be due to the number of grains in each of the simulations. Future work would include investigating the affect of grain number on the corresponding stress-strain relationship. It should be noted that the crustal composition in the accreted case has an impurity factor which is higher than the limits placed on the factor through neutron star cooling \\citep{horowitz09b}. A neutron star which has not undergone any accretion events is unlikely, and in the accreted crust case the impurity factor is higher than observational constrains, the results from these two cases may represent two extremes of what may be occurring in nature. In a low mass X-ray binary systems with an accretion rate of $10^{-9}\\mathrm{\\msun\\,yr^{-1}}$ the crust could be replaced via accretion in $10^7$ years \\citep{chamel08}. In the case of dim isolated neutron stars, as well as AXPs and SGRs the accretion rate has been estimated to range from $3.2\\times10^{14}\\mathrm{g\\,s^{-1}}$ to $4.2\\times10^{17}\\mathrm{g\\,s^{-1}}$ or in solar mass units the accretion rates are $5\\times10^{-12}\\mathrm{\\msun\\,yr^{-1}}$ to $7\\times10^{-9}\\mathrm{\\msun\\, yr^{-1}}$ \\citep{alpar01}. For an isolated neutron star these accretion rates correspond to a total crustal mass of $0.01\\msun$ \\citep{chamel08} being replaced in ${\\sim}10^{9}$ to ${\\sim}10^{6}$years. The accretion rate can also be estimated by attributing the observed luminosity to accretion. For at 10\\,km and 1.4\\msun dim isolated neutron star, with a luminosity of $10^{32}\\rm{erg}$, this corresponds to $\\dot{M} \\sim 10^{-14}\\msun\\,yr^{-1} $ and the crust being replaced by accretion in $10^{12}$ years. As such, the crustal composition calculated for  the non-accreting neutron stare, presented in this paper, would hold for those isolated neutron stars which have an accretion rate on the lower end of the rates.   The molecular dynamics simulations reported here leave many venues for future work  in order to fully understand the neutron star crust. The neutron star crust would not be expected to have a perfect lattice structure, thus it would be important to understand the effect of defects added to the simulation, as well how the orientation of the applied shear also affects the material. Fully characterizing the material as the simulation progresses would give an indication of how the structure of the crustal material changes with applied shear and cooling, this may have an effect on the material properties, as well as conductivity, of the crust. In all the simulations magnetic fields were not directly placed within the simulation, the addition of magnetic fields may have an effect on the mechanical properties, especially as with high magnetic fields there is a change in the electron screening \\citep{rishi11}. The simulation results reported here indicate upper limits on the crustal properties, further simulations may bring us closer to understanding the true nature of the neutron star crust."
    },
    "1208/1208.0367_arXiv.txt": {
        "abstract": "We present simulation results for the formation and long-term evolution of a primordial protostellar disk harbored by a first star. Using a 2+1D nonaxisymmetric thin disk numerical simulation, together with a barotropic relation for the gas, we are able to probe ${\\sim}20~\\mbox{kyr}$ of the disk's evolution. During this time period we observe fragmentation leading to loosely bound gaseous clumps within the disk. These are then torqued inward and accreted onto the growing protostar, giving rise to a burst phenomenon. The luminous feedback produced by this mechanism may have important consequences for the subsequent growth of the protostar. ",
        "introduction": "Although the existence of protostellar disks is a ubiquitous outcome of the present-day star formation process, the importance of these structures in the evolution of the first stars in the early universe has come to light only recently \\citep{stacy2010,clark2011}; with their significance not fully understood. Numerical simulations by several authors have converged upon the idea that collapsing primordial cores produce a much more structure-rich environment than previously thought: binary pairs \\citep{machida2008}, embedded cluster formation \\citep[e.g.,][]{clark2008}, and vigorous small-scale fragmentation \\citep{greif2012}, are all likely outcomes of primordial disk formation and fragmentation.  A major limitation of the aforementioned calculations however, is the inability to follow the disk evolution for much more than a thousand years. Herein we present 2+1D numerical simulations of the formation and evolution of a primordial circumstellar disk over time scales that are physically relevant to the global picture of Population III protostellar evolution. ",
        "conclusions": "We follow the evolution of isolated, gravitationally unstable, prestellar primordial cores as they collapse into the protostar and disk formation stage. The dynamics of the entire system are followed self-consistently on a single, globally defined computational domain. Here we focus on the innermost region where the disk forms around the central star. In Figure~\\ref{fig:surfacemassdensity} we present a map of the surface mass density from a reference model, together with profiles of the azimuthally averaged surface mass density and gravitational torque. Within a time $t~{\\leq}~5~\\mbox{kyr}$ the disk begins to experience the first episodes of fragmentation. Fragments are typically formed at radial distances of ${\\sim}100~\\mbox{AU}$ from the central protostar, and their subsequent orbital dynamics depend on the sign of the gravitational torque acting on them. In our simulations all fragments are seen to either migrate inward, or are tidally dispersed while still in the outer disk. These outcomes result from a combination of the gravitational torques from the passage of spiral arms and interactions with other fragments. The fragmentation process is driven by the nearly continuous supply of new material infalling from the envelope.  \\begin{figure}[t!] \\includegraphics[width=0.333\\textwidth]{./figure1.eps} \\includegraphics[width=0.333\\textwidth]{./figure2.eps} \\includegraphics[width=0.333\\textwidth]{./figure3.eps} \\caption{\\textbf{Left:} Surface mass density map ${\\sim}11~\\mbox{kyr}$ after the formation of the primordial protostar. The gradient indicates the logarithm of the surface mass density in $\\mbox{g}~\\mbox{cm}^{-2}$. \\textbf{Center:} Azimuthally averaged radial surface mass density profile. The dotted line shows a $r^{-1.5}$ profile for comparison. \\textbf{Right:} Azimuthally averaged radial profiles of the surface mass density (solid) and normalized gravitational torque (dashed).} \\label{fig:surfacemassdensity} \\end{figure}  \\begin{figure}[t!] \\includegraphics[width=0.333\\textwidth]{./figure4.eps} \\includegraphics[width=0.333\\textwidth]{./figure5.eps} \\caption{\\textbf{Left:} Mass accretion rates, from the disk onto the central protostar (solid), and from the outer envelope onto the disk (dashed). \\textbf{Right:} Accretion luminosity associated with the infall from the disk onto the central protostar.} \\label{fig:accluminosity} \\end{figure}  Fragments surviving the inward migration and transition through the inner domain wall are likely to be tidally destroyed as they plunge into the central star, releasing their gravitational energy in the form of strong luminosity outbursts (lasting approximately 100 years). Estimating the protostellar radius with simple power-law approximations that characterize the distinct phases of the star's internal dynamics \\citep{omukai2003,smith2011}, we calculate the luminosity associated with each burst of accretion (Figure~\\ref{fig:accluminosity}). The luminosity generated through this burst mode of accretion is as high as a few times $10^6~\\mbox{L}_{\\odot}$. That the predominant fraction of material accreted onto the protostar occurs via the burst mode suggests previous estimates of the ionizing UV-flux associated with the earliest phases of the protostar's evolution \\citep{mckee2008} may be underestimated by as much as a factor of $10$. Energy deposition into the environment of this order has the potential to significantly affect the nature of mass accretion onto primordial stars in the early universe.   \\begin{theacknowledgments} EIV aknowledges support from the RFBR grant 10-02-00278. \\end{theacknowledgments}"
    },
    "1208/1208.5890_arXiv.txt": {
        "abstract": "From the GOES-12/SXI data, we studied the initial stage of motion for six rapid (over 1500 km/s) \"halo\" coronal mass ejections (HCMEs) and traced the motion of these HCMEs within the SOHO/LASCO C2 and C3 field-of-view. For these HCMEs the time-dependent location, velocity and acceleration of their fronts were revealed. The conclusion was drawn that two types of CME exist depending on their velocity time profile. This profile depends on the properties of the active region where the ejection emerged. CMEs with equal ejection velocity time dependence originate form in the same active region. All the HCMEs studied represent loop-like structures either from the first moment of recording or a few minutes later. All the HCMEs under consideration start their translational motion prior to the associated X-ray flare onset. The main acceleration time (time to reach the highest velocity within the LASCO/C2 field-of-view) is close to the associated flare X-ray radiation intensity rise time. The results of (Zhang and Dere, 2006) on the existence of an inverse correlation between the acceleration amplitude and duration, and also on the equality of the measured HCME main acceleration duration and the associated flare soft X-ray intensity rise time are validated. We established some regularities in the temporal variation of the angular size, trajectory, front width and the HCME longitude-to-cross size ratio. ",
        "introduction": "\\label{S-Introduction}  Examining the properties of a coronal mass ejection (CME) at the initial stage of its motion is a necessary stage in the process of identifying the physical mechanisms involved in CME formation. Over the last 10 years, a considerable amount of work has been carried out in which CME kinematics was studied right after their formation using the data of various telescopes (see Gallagher et al., 2003; Zhang and Dere, 2006; Temmer el al., 2008; Mari\u010di\u0107 et al., 2009; Patsourakos et al., 2010; Temmer et al., 2010 and the quoted literature in these papers). These papers obtained important results regarding the initial stage of CME motion.  Zhang and Dere (Zhang and Dere, 2006) came to the conclusion that CME motion can be subdivided into three stages: \\textit{i\\textup{)}} initial stage, when CME speed slowly increases, \\textit{ii\\textup{)}} main acceleration phase lasting several minutes to several hours, and \\textit{iii\\textup{)}} quiet motion phase at an approximately constant speed. They showed that the post-acceleration CME speed and kinetic energy correlate with the maximum value of soft x-ray radiation intensity $I_{\\rm SXR}(t)$ from  the flare area relating to the CME (Moon et al., 2003; Burkepile et al., 2004; Vr\\v{s}nak et al., 2005). The flare and CME are considered to be related if the place and moment of their emergence are close. It has been established that the kinematics of many CMEs is synchronised with the $I_{\\rm SXR}(t)$ behaviour (Gallagher et al., 2003; Zhang and Dere, 2006; Mari\\v{c}i\\'{c}et al., 2007; Patsourakos et al., 2010). This manifests itself in the time profile of CME speed $V(t)$ in its main acceleration phase being close to the time dependence of the soft X-ray radiation (SXR) intensity, $I_{\\rm SXR}(t)$, from the CME-related flare area. There is an inverse correlation between the main acceleration of a CME and its measured duration, as well as between the former and the time it takes for $I_{\\rm SXR}(t)$ to increase, from the flare onset to the moment $I_{\\rm SXR}(t)$ reaches its maximum (Zhang and Dere, 2006; Mari\\v{c}i\\'{c} et al., 2009). Here the acceleration value was defined as the maximum ejection speed divided by the main acceleration time. It is concluded in (Temmer et al., 2008, 2010) that the time profile of the main CME acceleration a(t) is synchronised with the profile of  hard X-ray radiation intensity $I_{\\rm HX}(t)$ from Reuven Ramaty High-Energy Solar Spectroscopic Imager (RHESSI; Lin et al., 2002). According to (Patsourakos et al, 2010), the profile $a(t)$ is close to the time dependence of a derivative of the soft x-ray radiation intensity $dI_{\\rm SXR}/dt$. This parameter is sometimes used as an analogue of $I_{\\rm HX}(t)$. This is due to the existence of the  Neupert effect (Neupert, 1968), according to which the temporal changes of $I_{\\rm HX}(t)$ or of the microwave radiation intensity during the pulse phase of a flare are close to the time profile $dI_{\\rm SXR}/dt$.  CMEs for which the initial motion stage was studied and their kinematic characteristics were compared with soft and hard x-ray radiation intensity were mostly limb CMEs. Their sources are relatively close to the solar limb, and the axes of such ejections are presumably located near the sky plane. The 3-D geometrical and kinematic characteristics of such CMEs are believed to be close to those determined from sky-plane observations of these ejections, i.e. for limb CMEs, the projective effects do not make a large impact on determining their real parameter values.  A special group is distinguishable among all observable coronal mass ejections, called \"halo\" CMEs (HCMEs). They are observable in the field of view of a coronagraph as areas of enhanced brightness completely surrounding the occulting disk and expanding in all directions (Howard et al., 1982). Some HCMEs move towards the ground-based observer (frontside HCMEs), while others move in the opposite direction (backside HCMEs). In the former case, the CME sources are on the visible solar disk. It is such HCMEs that play a significant role in space weather: their influence on the Earth magnetosphere can lead to strongest geomagnetic storms (Gopalswamy, 2009).  HCMEs are quite convenient objects for exploring the formation mechanisms for coronal mass ejections and the initial phase of their motion, especially when HCME sources are relatively not far from the center of the visible solar disk. This is due to the fact that such cases provide the opportunity to \"see\" all the phenomena at the site of CME formation and in the adjoining areas of the solar disk, as well as for a more correct examination of magnetic field dynamics in this area. The above merits compensate for the basic disadvantage that is inherent in studies of the  initial stage of HCME motion - increased influence of projective effects when determining the position, speed and acceleration of the leading edge (LE) of the HCME in three-dimensional space. The kinematic characteristics for the initial stage of two moving HCMEs as well as the relationship between these characteristics and the parameters of hard X-ray radiation from the HCME-related flare area have been discussed in, e.g. (Temmer et al., 2008).  Of special interest are CMEs moving at high speed in the field of view of the Large Angle and Spectrometric Coronagraph (LASCO; Brueckner et al., 1995) onboard the Solar and Heliospheric Observatory (SOHO), and featuring a short main acceleration phase varying from a few minutes to several tens of minutes. Fast CMEs are often found to be related to to powerful flares of the M and X (X-ray) classes.  This paper inspects the properties of six HCMEs selected from a group of the fastest-moving ejections, $V>1500$ km/s, observed over solar cycle 23. The properties include the kinematics, angular sizes and trajectories of the ejections. The connection between these HCMEs and solar flares has also been examined. The initial stage of HCME motion was identified based on high temporal resolution data of the Solar X-ray Imager (SXI; Hill et al. 2005) onboard the GOES-12 space observatory. This paper relies on the SXI as an instrument with a minimum time cadence of 60 seconds and spatial resolution of 5 arcsec per pixel.  One of the reasons for selecting the fastest HCMEs is that such ejections are, on the average, characterized by higher brightness than slow HCMEs both in the ejection body area and in the shock wave area. The same property is also valid for fast limb CMEs (Fainshtein, 2007). This allows for a more precise identification of the HCME body boundary and of the ejection-related shock front as well as for tracing the ejection over large distances. Moreover, the fast ejections are more often related to powerful x-ray flares and eruptive filaments. Finally, the fast HCMEs belong to more geoeffective events than do the slow HCMEs. ",
        "conclusions": "% \\label{S-Conclusion} We investigated the laws of the initial stage of movement of six halo coronal mass ejections using the data of several telescopes, and, first of all, GOES-12/SXI, SOHO/LASCO, SOHO/EIT. For the analysis, these HCMEs were selected from the group of the fastest ejections (with $V>1500~km/s$) connected with powerful flares. The choice of such events was determined, first, by the role of fast HCME in space weather, and second, the opportunity to reveal them at the stage of their formation. The HCME events selected for analysis took place place in 2003 and 2005. We used SXI as the telescope with high time cadence, though during the period when the HCME events analysed took place the telescope TRACE with higher time cadence was working. But the use of TRACE images for studying the initial stage of the movement of all selected ejections turned out to be impossible because of the limited field of view of this telescope and some features of its functioning.  We were not able to find eruptive filaments connected with ejection in any of the HCME considered on the $H_{\\rm \\alpha}$ images. At the same time, we established that, before the occurrence of ejection in five of the six HCME events (except the HCME event on 29 Oct 2003), the loop-like structure identified as the beginning of HCME using SXI data, is observed as a single emission coronal loop in the extreme ultraviolet spectral line at $\\lambda =19.5~nm$ on the EIT data image. All the loop-like structures become more active and begin slow forward movement before the beginning of the flare connected with the HCME. In this work we do not discuss questions concerning the reasons for formation and activation of these emission loop-like structures. It is possible that it is the flux-ropes existing in the active area for some time and becoming more active as a result of the infringement of magnetic field balance. Possible mechanisms of such balance infringement are discussed in several works [Forbes et al., 2006; Howard, 2011]. The occurrence of an HCME event on 29 Oct 2003 was apparently connected with the movement not of a separate loop, as in the other events studied, but with the movement of an arcade of loops or groups of unconnected loops.  We have shown that the time dependence of the $V(t)$ speed profile of the fast HCME can be of two types. So, for the HCME event on 17 Jan 2005, the $V(t)$ profile quickly reaches maximum value before entering the field of view of LASCO C2, then it falls sharply over a short period of time and then slowly decreases (see Figure 2\u0410). The two HCME observed on 15 Jan 2005 in this active region have close kinematic properties. The HCMEs consistently observed on 22 and 23 August 2005 close to the limb from the same active region (see Figure 4\u0410) have another $V(t)$ profile change character. For these HCME events the $V(t)$ profile quickly increases before the appearance of the ejection into the field of view of LASCO C2 coronagraph and then it continues to increase slowly and reaches maximum value at a large distance from the Sun's surface. It turned out that the HCMEs arising in one and the same active region have identical time speed profiles. We  assume that HCMEs with such different time speed profiles relate to two different classes of coronal mass ejections. It was ascertained that HCMEs of the first type are formed in active regions with a complex configuration of sunspots and magnetic field and with a large sunspot area, and HCMEs of the second class are formed in active regions with more simple sunspot and magnetic field configurations.  As in a number of previous papers by other authors we came to the conclusion that the kinematics of the HCMEs analysed is synchronized with the time variation of intensity of soft X-ray radiation $I_{\\rm SXR}(t)$ from the flare area connected with the ejection [Gallagher et al., 2003; Zhang and Dere, 2006; Mari\\v{c}i\\'{c} et al., 2007; Patsourakos et al., 2010]. It turned out, that the $t_{\\rm ACC}$ time of the basic acceleration of all HCMEs investigated is close to the time of increase $I_{\\rm SXR}(t)$ from the flare's beginning till the moment of $I_{\\rm SXR}$ maximum value. The HCME maximal measured acceleration adopts the value $a_{\\rm MAX}\\approx 0.9-10~km/s^2$ and is close to the value obtained by dividing the HCME maximal velocity $V_{\\rm MAX}$ by the X-ray intensity rise time $t_{\\rm SXR}$. It is shown that there is an inverse correlation between $V_{\\rm MAX}/t_{\\rm ACC}$ and $V_{\\rm MAX}/t_{\\rm SXR}$ on the one hand and $t_{\\rm ACC}$ and $t_{\\rm SXR}$ - on the other. This is consistent with the results in [Zhang and Dere, 2006]. At the same time, this dependence in our study was obtained for the values $V_{\\rm MAX}$ and $\u0430_{\\rm MAX}$ that are much higher than in [Zhang and Dere, 2006]. Acceleration $a(t)$ of HCMEs is closely connected with the intensity of rigid X-ray radiation $I_{\\rm HX}(t)$ from the flare area. But conformity between $a(t)$ and $I_{\\rm HX}(t)$ values varies from event to event.  At the initial movement stage, the trajectories of several HCMEs investigated are curvilinear and deviate from the equator. At the same time, the HCMEs originating in the Sun's northern hemisphere deviate towards the North Pole, whereas the HCMEs originating in the southern hemisphere deviate towards the South Pole.  We studied the time variation in the HCME angular size. The angular size of all the HCMEs under consideration is shown to augment with time from their translational motion onset, and may increase by a factor of $\\sim 4.5$ within the SXI field-of-view. The characteristic time for the increase in angular size by a factor of 1.5 is 5-10 minutes. The initial angular size of four HCMEs on SXI data images does not exceed $10^\\circ$, and for one ejection - $14^\\circ$. For the HCME event on 29 Oct 2003 there is uncertainty as to the estimation of the initial angular size of the ejection since it is formed from three small loops or from a loop arcade. After the ejection was formed its initial angular size was about $32^\\circ$.  The change  in the time variation ratio $d_{\\rm H}/d_{\\rm W}(t)$ of the HCME longitudinal size and its cross-section size was investigated. We showed that for each HCME under consideration its longitude-to-cross size ratio increases with time within the first minutes of motion, and then the character of change in this parameter turns out different for each ejection depending on time. In some cases the longitudinal size of the HCME changes quicker than the crosswise, in others vice versa. This ratio turns out to be close to unity within the LASCO field-of-view for three HCMEs, and does not vary for some time. This may reflect the HCME transition to the self-similar expansion mode. But taking into account that we are investigating the movement of halo coronal mass ejections such an inter-pretation $d_{\\rm H}/d_{\\rm W}(t)$ in the field of view of LASCO coronagraphs can appear incorrect.  We also investigated the time variation of HCME's front width $\\Delta f(t)$. The definition of the HCME's front width is shown in Figure 1(E) where the HCME's front is limited by the top and the bottom plusses (and see also the illustration on Figure 15(A)). It has been established, that the HCME's front width for five of the six events of HCMEs studied increases non-monotonously with time. In one case the tendency to increase of $\\Delta f(t)$ arises at the end of the time period for which it was possible to measure this parameter (see Figure 15).  \\begin{acks} We established that the front width for all the HCMEs under consideration augments non-monotonously with time in general case. The authors are grateful to the GOES and GOES-12/SXI, RHESSI, TRACE, SOHO/EIT, SOHO/LASCO and MarkIV (MLSO) teams for the possibility to freely use the data from these instruments. The authors appreciate V.V. Grechnev's most beneficial discussions of this paper. \\end{acks}"
    },
    "1208/1208.2965_arXiv.txt": {
        "abstract": "Relativistic outflows in the form of jets are common in many astrophysical objects. By their very nature, jets have angle dependent velocity profiles, $\\Gamma = \\Gamma(r,\\theta,\\phi)$, where $\\Gamma$ is the outflow Lorentz factor. In this work we consider photospheric emission from non-dissipative jets with various Lorentz factor profiles, of the approximate form $\\Gamma \\approx \\Gamma_{\\rm 0}/[(\\theta/\\theta_{\\rm j})^p + 1]$, where $\\theta_{\\rm j}$ is the characteristic jet opening angle. In collimated jets, the observed spectrum depends on the viewing angle, $\\theta_{\\rm v}$. We show that for narrow jets ($\\theta_{\\rm j} \\Gamma_{\\rm 0} \\lesssim few$), the obtained low energy photon index is $\\alpha \\approx -1$ ($dN/dE \\propto E^\\alpha$), independent of viewing angle, and weakly dependent on the Lorentz factor gradient ($p$). A similar result is obtained for wider jets observed at $\\theta_{\\rm v} \\approx \\theta_{\\rm j}$. This result is surprisingly similar to the average low energy photon index seen in gamma-ray bursts. For wide jets ($\\theta_{\\rm j} \\Gamma_{\\rm 0} \\gtrsim few$) observed at $\\theta_{\\rm v} \\ll \\theta_{\\rm j}$, a multicolor blackbody spectrum is obtained. We discuss the consequences of this theory on our understanding of the prompt emission in gamma-ray bursts. ",
        "introduction": "\\label{sect:introduction}  Photospheric emission from highly relativistic outflows was early considered as an explanation for prompt gamma-ray bursts (GRBs, \\citealt{Goo:1986, Pac:1986}). It is a natural consequence of the fireball model, where the optical depth at the base of the outflow is much larger than unity (e.g. \\citealt{Pir:1999}). Moreover, photospheric emission provides a natural explanation to the small dispersion of the sub-MeV peak and to the high prompt emission efficiency observed \\citep{MesRee:2000}. However, the observed spectrum usually appears significantly broader than a Planck spectrum \\citep{PreEtAl:2000}, being well fitted by a smoothly broken power law (the Band function, \\citealt{BanEtAl:1993}). Thus, the prompt emission has commonly been associated with synchrotron emission originating from kinetic energy dissipation outside the photosphere \\citep{ReeMes:1994}. However, in recent years it has become clear that optically thin synchrotron emission is incompatible with the hard low energy slopes observed in a substantial fraction of GRBs \\citep{PreEtAl:1998, KanEtAl:2006, Bel:2012}. This has raised the need for alternative ideas.  One appealing idea is broadening of the thermal spectrum emitted from the photosphere. The emerging photon spectrum from a static, optically thick, relativistic outflow can be widened in two ways:  \\begin{enumerate} \\item Energy dissipation below the photosphere can heat electrons above the equillibrium temperature. These electrons emit synchrotron emission and Comptonize the thermal photons, thereby modifying the Planck spectrum \\citep{ReeMes:2005, PeeMesRee:2005, PeeMesRee:2006}. The dissipation can be caused by shocks \\citep{ReeMes:2005, LazMorBeg:2009, RydEtAl:2011}, dissipation of magnetic energy \\citep{Tho:1994, SprDaiDre:2001, GiaSpr:2005, ZhaYan:2011} or collisional processes \\citep{Bel:2010, VurEtAl:2011}. \\item The photospheric radius is angle dependent \\citep{AbrNovPac:1991, Pee:2008}. Moreover, it was shown by \\citet{Pee:2008} that photons make their last scatterings at a distribution of radii and angles. The observer sees simultaneously photons emitted from a large range of radii and angles. Therefore, the observed spectrum is a superposition of comoving spectra. The Doppler boost is a function of angle, and the comoving temperature decreases with radius through adiabatic cooling. Depending on the outflow properties, this geometrical broadening can form observed spectra which appears significantly different from the Planck spectrum. \\end{enumerate}  Photospheric emission in the context of spherically symmetric outflows has been studied by several authors. \\citet{Goo:1986} considered a highly relativistic outflow in the context of cosmological GRBs. It was realized that the observed spectrum is broader than blackbody, however the analysis was one-dimensional. \\citet{AbrNovPac:1991} realized that the two-dimensional shape of the photosphere in a relativistic, spherically symmetric wind is in fact concave and symmetric around the line-of-sight (LOS). This can be understood as a consequence of the dependence on viewing angle of the optical depth of relativistically moving media. \\citet{Pee:2008} found a simple expression for the photospheric radius, $R_{\\rm ph} \\propto (\\theta^2/3 + 1/\\Gamma^2)$, where $\\theta$ is the angle measured from the LOS, $\\Gamma \\equiv (1 - \\beta)^{-1/2}$ is the outflow bulk Lorentz factor and $\\beta$ is the outflow speed in units of the speed of light. \\citet{Pee:2008} extended the photospheric emission model by recognizing the importance of considering photons from the entire emitting volume, introducing probability density distributions for the last scattering photon positions. \\citet{Bel:2011} took a different approach, solving the radiative transfer equation in the relativistic limit. The approximate probability densities used by \\citet{Pee:2008} was validated by \\citet{Bel:2011}.  All above mentioned works considered spherical explosions. As we show below, this is a good approximation for collimated outflows as long as the characteristic jet opening angle is much larger than $1/\\Gamma$ and the outflow is observed at viewing angles much smaller than the jet opening angle. Within the collapsar model \\citep{MacWoo:1999} the jet is collimated by the pressure of the surrounding gas as it drills its way through the collapsing progenitor star. In such jets the position of the observer relative to the jet axis can affect the observed spectrum. Here we develop the theory of photospheric emission in collimated outflows, calculating the expression for the observed spectrum at any viewing angle.  The mechanism responsible for jet collimation is not fully understood. Hydrodynamic simulations of GRB jets after the launching phase (e.g., \\citealt{ZhaWooMac:2003, MorLazBeg:2007, MizNagAoi:2011}) show the jet drilling through the stellar envelope, pushing material towards the sides, forming a hot cocoon (e.g. \\citealt{AloEtAl:2000}). The surrounding gas then acts to collimate the jet. \\citet{ZhaWooMac:2003} extracted angular profiles of the local rest mass density, total energy flux and outflow Lorentz factor at certain radii from the simulations. The resulting profiles can in many cases be approximated as constant within a characteristic jet opening angle, then decreasing as power laws towards the jet edge. In this work we adopt a similar parametrization of the angular profile of the bulk jet Lorentz factor, with a characteristic jet opening angle, power law index and normalization as free profile parameters.  We develop a model for photon propagation in the context of a steady, optically thick, axisymmetric jet with angle dependent electron number density, photon number density and bulk Lorentz factor. We compute the observed spectrum taking into account contributions from the entire emitting volume as seen by an observer located at any viewing angle. As an example solution, we consider fireball dynamics as a way to relate the angle dependent parameters, in combination with the assumed angular Lorentz factor profile. We developed a Monte Carlo simulation, unique in its ability to calculate photon propagation in a non-spherical explosion. We use this simulation to analyze the importance of angular bulk photon diffusion.  We show that for a large region in the parameter space, the low energy photon index is approximately $\\alpha \\approx -1$ (where $dN/dE \\propto E^\\alpha$). This is similar to the average value observed in GRBs \\citep{KanEtAl:2006, NavEtAl:2011, GolEtAl:2012}. Furthermore, we present analytical expressions for the important energies and photon indices of the observed spectrum as functions of the free model parameters. We show that photospheric emission by itself can account for observations of GRB spectra below the peak energy without the need for energy dissipation below the photosphere or additional radiative processes, provided that the characteristic jet opening angle is not much larger than $\\sim few/\\Gamma$ or if the outflow is viewed off-axis. Although we focus on GRB parameters, the results are general and can readily be applied to any optically thick, relativistic outflow with lateral outflow properties such as active galactic nuclei.  This paper is organized as follows. In \\S \\ref{sect:model} we develop the analytical expression for the observed spectrum. We qualitatively explain the spectral features in \\S \\ref{sect:the observed spectrum} in terms of contributions from different jet regions. The Monte Carlo code is explained in \\S \\ref{sect:numerical simulations}, and in \\S \\ref{sect:results} we present both simulated and numerically integrated spectra for outflow parameters relevant to GRBs. We discuss model sensitivity and photon time delays in \\S \\ref{sect:discussion} and summarize our results in \\S \\ref{sect:summary and conclusions}. ",
        "conclusions": "\\label{sect:summary and conclusions}  In this work we develop the theory of photospheric emission from relativistic jets with angle dependent outflow properties. We consider a three parameter angular Lorentz factor profile (Eq. \\ref{eq:gamma profile}) where the Lorentz factor is approximately constant, $\\Gamma = \\Gamma_{\\rm 0}$ within a characteristic jet opening angle $\\theta_{\\rm j}$ and then decreasing approximately as $\\Gamma \\propto \\theta^{-p}$ towards the outer jet edge. The shape of the profile is motivated by the results of hydrodynamical simulations of jet propagation through the progenitor star (e.g. \\citealt{AloEtAl:2000, ZhaWooMac:2003, MorLazBeg:2007, MizNagAoi:2011}). In \\S \\ref{sect:model} the expression for the observed spectrum is obtained analytically by integrating the local emissivity over all radii and angles (Eq. \\ref{eq:long spectrum}). We derive approximate analytical expressions for the important spectral features in \\S \\ref{sect:the observed spectrum}. Comparing jetted outflows to spherical outflows, we show that softening of the spectrum below the peak energy is expected for an on-axis observer. This is a consequence of weaker radial beaming of photons at high latitudes due to the decreasing Lorentz factor. A Monte Carlo simulation was developed to investigate the importance of full photon propagation below the photosphere (\\S \\ref{sect:numerical simulations}). In \\S \\ref{sect:results} we present spectra obtained by numerical integration of Eq. \\ref{eq:long spectrum} as well as the Monte Carlo simulations for different profile parameters and viewing angles (Figures \\ref{fig:g100p1j001_vary_v} to \\ref{fig:g100v0j001_vary_p}).  The most important result of this paper is that { the photospheric spectrum below the thermal peak may be significantly softer than blackbody, as a consequence of geometrical broadening. In particular,} we obtain a photon index $\\alpha \\approx -1$ below the peak energy for a large region of the considered parameter space. For narrow jets ($\\theta_{\\rm j} \\lesssim few/\\Gamma_{\\rm 0}$) with Lorentz factor gradients $1 \\leq p \\leq 4$ observed at any viewing angle, we find $-1 \\gtrsim \\alpha \\gtrsim -0.5$. For jets with $\\theta_{\\rm j} \\approx 1/\\Gamma_{\\rm 0}$ and $p \\geq 1$, $\\alpha \\approx -(1/4)(1 + 3/p)$ (see Eq. \\ref{eq:spectrum solution 4}). Observing wider jets ($\\theta_{\\rm j} \\gtrsim 5/\\Gamma_{\\rm 0}$) at $\\theta_{\\rm v} \\approx \\theta_{\\rm j}$ results in similar soft spectra with $\\alpha \\approx -1$. { However, spectra from wide jets observed at small angles ($\\theta_{\\rm v} \\ll \\theta_{\\rm j}$) appears similar to the spectrum from a spherical wind (i.e. close to blackbody but with $\\alpha \\approx 0.4$ at $E = 10^{-2} \\, E_{\\rm peak}$) for several decades below the peak energy. This may explain the hard low energy photon indices observed in some GRBs ($\\alpha \\approx 0$ \\citep{GolEtAl:2012}, see further discussion below).} However, observing the outflow at viewing angles $\\theta_{\\rm v} \\approx 0$ is less likely than $\\theta_{\\rm v} \\approx \\theta_{\\rm j}$. Additionally, increasing the viewing angle causes the observed peak energy to decrease. The decrease is slower for low Lorentz factor gradients. Therefore, such outflows may still be observed at $\\theta_{\\rm v} \\approx few \\times \\theta_{\\rm j}$.  Photon diffusion primarily towards regions of higher Lorentz factor is observed in the simulations. For outflows with large Lorentz factor gradient ($p \\gtrsim 4$) this bulk propagation of photons has to be taken into account when computing the observed spectrum. Comptonization of the photons that make repeated scatterings towards regions of larger Lorentz factor can produce photons with energies significantly higher than the local temperature. Evidence for this can be seen in Figure \\ref{fig:g100v0j001_vary_p} ($p = 4$), where an approximate high energy power law is formed above the peak energy. Photon propagation in plasma with a steep Lorentz factor gradient will be further explored in future works.  The observed photospheric spectrum from a spherical outflow has previously been considered in the literature (e.g. \\citealt{Bel:2010, PeeRyd:2011}). In the limiting case of outflows with $\\theta_{\\rm j} \\gtrsim 5/\\Gamma_{\\rm 0}$ observed at $\\theta_{\\rm v} \\ll \\theta_{\\rm j}$, we obtain similar results. For larger viewing angles or smaller jet opening angles, geometrical broadening of the spectrum has to be considered.  Although we consider a static source, as shown in \\S \\ref{subsect:time delays} there is a time delay associated with high-latitude photons. The time delay increases with decreasing observed photon energy. The time delay at a specific energy is longer for more narrowly collimated jets. For a narrow jet ($\\theta_{\\rm j} = 1/\\Gamma_{\\rm 0}$) with typical outflow parameters characterizing GRBs, the time delay is $\\approx 1 \\, {\\rm s}$ at $E \\approx 10^{-2} E_{\\rm peak}$. Spectral analysis of prompt GRB emission using smaller time bins than this time delay may reveal harder spectra within the energy range $10^{-2} E_{\\rm peak} \\lesssim E \\lesssim E_{\\rm peak}$ than what is predicted by the static model.  For this work we have considered outflows with angle independent luminosity, and so the shape of the Lorentz factor profile is fully determined by the angle dependent baryon loading. As shown above, the spectral slope in the range $10^{-2} E_{\\rm peak} \\lesssim E \\lesssim E_{\\rm peak}$ is formed by photons making their last scattering at $\\theta \\lesssim 5/\\Gamma_{\\rm 0}$. Therefore, this assumption is a good approximation for jets where $dL/d\\Omega \\approx const$ for $\\theta \\lesssim 5/\\Gamma_{\\rm 0}$. This requirement is fulfilled for model JA in \\citet{ZhaWooMac:2003} close to the largest radius of the simulation ($r = 2.1 \\times 10^{10} \\, {\\rm cm}$). In our notation, $\\Gamma_{\\rm 0} \\approx 140$, $\\theta_{\\rm j} \\approx 0.017 \\approx 2.5/\\Gamma_{\\rm 0}$ and $p \\approx 6.5$ for model JA as estimated from Figures 8 and 9 in \\citet{ZhaWooMac:2003}.  Furthermore, we consider non-dissipative fireball dynamics. The dynamics of the outflow are dependent on the dominant form of energy carried by the jet as well as energy dissipation. In particular, the radial scalings of the comoving temperature and Lorentz factor in magnetically dominated outflows are expected to differ significantly from the scalings of thermal fireballs. \\citet{Gia:2012} and \\citet{Bel:2012} considered the scaling of $E_{\\rm peak}$ in dissipative outflows. Heating keeps the comoving temperature approximately constant in the range $R_{\\rm ph}/30 \\lesssim r \\lesssim R_{\\rm ph}$, in contrast to the non-dissipative outflows considered here. However, the framework presented in \\S \\ref{sect:model} is general and may be applied to any relativistic, optically thick outflow.  Since the causally connected parts of the outflow are separated by angles $\\approx 1/\\Gamma$, one may consider outflows with Lorentz factor variations at angular scales of $\\Delta\\theta \\approx few/\\Gamma$. Geometrical broadening of the observed spectrum is expected as a consequence of the beaming of photons being a function of angle from the jet axis, in much the same way as for the jets considered in this work. The spectral slope below the peak energy is expected to depend on both the typical angular scale as well as the amplitude of the Lorentz factor variations.  Spectral broadening by energy dissipation in regions of moderate optical depth may be combined with geometrical broadening. As the observed spectrum below the peak energy is a superposition of comoving spectra, it is not sensitive to the exact shape of the comoving spectrum. Comptonization of the comoving spectrum by electrons which are heated by energy dissipation (e.g. \\citealt{Bel:2010, LazBeg:2010}) can shape the spectrum above the peak energy, while geometrical broadening forms the spectrum below the peak energy.  The observed low energy photon index varies between bursts, forming an approximately Gaussian distribution centered at $\\alpha \\approx -1$ with a full width at half maximum of $\\approx 1$ (e.g. \\citealt{GolEtAl:2012}). Within the framework presented in this paper, the distribution could naturally be interpreted as a result of viewing angle variations. { In particular, observing a wide jet at zero viewing angle results in $\\alpha \\approx 0.4$.} The exact $\\alpha$-distribution predicted by our model is hard to obtain, since it depends on detector characteristics. However, a clear prediction of the model is that of softening of the low energy photon index with increasing viewing angle for jets with $\\theta_{\\rm j} \\gtrsim few/\\Gamma_{\\rm 0}$.  GRB 090902B is a burst of special interest, due to its unusual spectral shape \\citep{AbdEtAl:2009}. The prompt spectrum consists of a sharply peaked Band component along with a wide power law. The hard low energy photon index and the narrow width of the Band component poses an extreme challenge for optically thin emission models. Therefore, a photospheric origin of the Band component in this burst seems inevitable \\citep{RydEtAl:2010, RydEtAl:2011, ZhaEtAl:2011, PeeEtAl:2012}. Since the Band component appears close to blackbody, processes that are expected to modify the photospheric spectrum can be constrained for GRB 090902B. The rate of energy dissipation in regions of moderate optical depth must be relatively low as the observed Band component is not severely distortened from the Planck spectrum. Within the framework considered in this paper, further constraints can be set. The requirement for geometrical broadening to not significantly distort the observed spectrum in the energy range $10^{-2} E_{\\rm peak} \\lesssim E \\lesssim E_{\\rm peak}$ is $\\theta_{\\rm v} \\ll \\theta_{\\rm j}$ and $\\theta_{\\rm j} \\gtrsim 5/\\Gamma_{\\rm 0}$. Depending on typical GRB outflow characteristics, the probability for such parameter combinations may be low. This would help explaining the rarity of similar events.  Multiple spectral components have been clearly identified in several GRBs after the launch of the {\\it Fermi} telescope \\citep{AbdEtAl:2009, AckEtAl:2010, AckEtAl:2011}. This was expected from a theoretical point of view, simply because of the wide energy range of the observations. Several models predict non-thermal spectral components as a result of kinetic or magnetic energy dissipation above the photosphere (e.g. \\citealt{MesEtAl:2002, ZhaYan:2011}). If dissipation occurs at $r \\gg R_{\\rm ph}$ the thermal and non-thermal components may be fitted separately (such as in 090902B \\citep{PeeEtAl:2012}). For dissipation at $r \\gtrsim R_{\\rm ph}$ the separation of thermal and non-thermal components may be less clear because of the coupling of the thermal photon field to the accelerated non-thermal electrons. In this work we consider the thermal component in isolation in order to demonstrate the geometrical broadening of the spectrum. However, a complete theory for the prompt GRB emission must explain all observed spectral components.  An outstanding problem for the photospheric interpretation of prompt GRB emission is the non-thermal spectra commonly observed. In particular, the average low energy photon index is not compatible with the Rayleigh-Jeans index ($\\alpha \\approx -1$ as compared to $\\alpha = 1$). In this work we show that either a narrow jet ($\\theta_{\\rm j} \\lesssim few / \\Gamma_{\\rm 0}$) with a moderate Lorentz factor gradient ($1 \\leq p \\leq 4$) observed at any viewing angle or a wide jet observed at $\\theta_{\\rm v} \\approx \\theta_{\\rm j}$ can naturally produce $-1 \\lesssim \\alpha \\lesssim -0.5$ through geometrical broadening, without a need for additional emission processes such as synchrotron emission to supply photons below the thermal peak."
    },
    "1208/1208.2154_arXiv.txt": {
        "abstract": "The \\Offline software framework for data analysis of the Pierre Auger Observatory is a set of computational tools developed to cater to the needs of a large and geographically dispersed collaboration established to measure the spectrum, arrival directions, and composition of ultra-high energy cosmic rays over a period of 20 years. One of its design goals was to facilitate the collaborative effort by allowing collaborators to progressively contribute small portions of code.  The observatory has grown over time and it has undergone improvements and additions that have tested the flexibility of the framework. The framework was originally thought to accommodate a hybrid view of cosmic ray detection, made of a surface and a fluorescence detector. Since then, the framework has been extended to include a radio antenna array and both under-ground and above-ground scintillator arrays.  Different tools from the framework have been used by other collaborations, notably NA61/Shine and HAWC. All these experiences accumulated over the years allow us to draw conclusions in terms of the successes and failures of the original design. ",
        "introduction": "The \\Offline software framework of the Pierre Auger Observatory \\cite{Argiro:2007qg, offline-chep07} provides tools and infrastructure to analyze data gathered by the observatory. The observatory is designed to measure the extensive air showers produced by the highest energy cosmic rays ($>$ 10$^{18.5}$ eV) with the goal of discovering their origins and shedding light on their composition. Two different techniques are used to detect air showers. First, a collection of telescopes is used to detect the fluorescence light produced by excited atmospheric nitrogen as the cascade of particles develops and deposits energy in the atmosphere. This method can be used only when the sky is moonless and dark, and thus has roughly a 15\\% duty cycle.  Second, an array of detectors on the ground is used to sample particle densities and arrival times as the air shower impinges upon the Earth\u2019s surface. Each surface detector consists of a tank containing 12 tons of purified water, instrumented with photomultiplier tubes to detect the Cherenkov light produced by passing particles. The surface detector has nearly a 100\\% duty cycle. A sub-sample of air showers detected by both instruments, called hybrid events, are very precisely measured and provide an invaluable tool for cross checks and energy calibration. The observatory, located in Mendoza, Argentina, was completed in 2008. It comprises 24 fluorescence telescopes overlooking an area instrumented with more than 1600 surface detectors spaced 1.5 km apart on a hexagonal grid.  The requirements of such a collaboration of over 400 scientists, from 17 countries, taking data over decades, imposes demands on the analysis software. It must be flexible enough in order to aggregate individual developments and allow the comparison of algorithms. It is essential that all physics code be \u201cexposed\u201d in the sense that any collaboration member must be able to replace existing algorithms with his or her own in a straightforward manner. This is meant to encourage independent analysis and ease the comparison of results. Finally, while the underlying framework itself may exploit the full power of C++ and object-oriented design, the portions of the code directly used by physicists should not assume a particularly detailed knowledge of these topics.  The framework was originally thought to handle simulation and reconstruction of events detected with the surface detector, the fluorescence detector or both, as well as simulation of calibration techniques and other ancillary tasks such as data preprocessing. It is essential that the software be extensible to accommodate future upgrades to the observatory instrumentation. Examples of upgrades that have been successfully included in the framework are AMIGA \\cite{amiga_2011}, HEAT \\cite{heat}, and the Auger Engineering Radio Array \\cite{Abreu:2011fb, aera_2011}.  The framework includes tools to facilitate multi-format file handling, and to provide access to event as well as time-dependent detector information which can reside in various data sources. A number of utilities are also provided, including a geometry package which allows manipulation of abstract geometrical objects independent of the choice of coordinate system. The distribution system incorporates unit and acceptance testing in order to support rapid development of both the core framework and contributed user code. The \\Offline framework can be obtained upon request and is released under BSD license. ",
        "conclusions": ""
    },
    "1208/1208.0826_arXiv.txt": {
        "abstract": "Among the most remarkable features of the stellar population of R136, the central, young, massive star cluster in the 30 Doradus complex of the Large Magellanic Cloud, are the single stars whose masses substantially exceed the canonical stellar upper mass limit of $150\\Ms$. A recent study by us, \\viz, that of Banerjee, Kroupa \\& Oh (2012; hereafter Paper I) indicates that such ``super-canonical'' (hereafter SC) stars can be formed out of a dense stellar population with a canonical initial mass function (IMF) through dynamically induced mergers of the most massive binaries. The above study consists of realistic N-body computations of fully mass-segregated star clusters mimicking R136 in which all the massive stars are in primordial binaries. In the present work, we study the formation of SC stars in the computed R136 models of Paper I in detail. Taking into consideration that extraneous SC stars form in the computed models of Paper I due to the primordial binaries' initial eccentricities, we compute additional models where all primordial binaries are initially circular. We also take into account the evolution of the mass of the SC stars and the resulting lifetime in their SC phase using detailed stellar evolutionary models over the SC mass range that incorporate updated treatments of the stellar winds. In all these computations, we find that SC stars begin to form via dynamical mergers of massive binaries from $\\approx 1$ Myr cluster age. We obtain SC stars with initial masses up to $\\approx250\\Ms$ from these computations. Multiple SC stars are found to remain bound to the cluster simultaneously within a SC-lifetime. However, we also note that SC stars can be formed at runaway velocities which escape the cluster at birth. These properties of the dynamically formed SC stars, as obtained from our computations, are consistent with the observed SC stellar population in R136. In fact, the evolutionary models of SC stars imply that had they formed primordially along with the rest of the R136 cluster, \\ie, violating the canonical upper limit, they would have evolved below the canonical $150\\Ms$ limit by $\\approx 3$ Myr, the likely age of R136. Thus according to the new stellar evolutionary models, primordially-formed SC stars should not be observable at the present time in R136. This strongly supports the dynamical formation scenario of the observed SC stars in R136. ",
        "introduction": "The study of the functional form of the number distribution of stars in galaxies, with which they are born, or the stellar initial mass function (IMF) has always been of fundamental importance \\citep{bast2010,pk2011}. The high-mass end of the stellar IMF is in particular focus due to the feedback it gives to the star-forming gas in the form of radiation pressure, kinetic energy from the stellar wind and supernova ejecta and as well the chemical enrichment from the latter source. An upper limit of $\\approx 60\\Ms$ to the mass of a star could be set by the ``Eddington limit'' \\citep{edd26} or the point of balance of gravity by the radiation pressure of a star. However, in practice stellar masses can easily exceed the Eddington limit since massive stars are not fully radiative but contain convective cores \\citep{kw90}. The next upper limit is set by the possibility of the destruction of a star due to thermal pulsations \\citep{sh59,bm94}. \\citet{stroh92} determined this limit to be at $\\imfmax\\approx 120-150\\Ms$ for [Fe/H]$\\approx 0$ and $\\imfmax\\approx 90\\Ms$ for [Fe/H]$\\approx -1$.  Difficulties also arise if one considers the growth of a proto-star via gas accretion which is crucial for massive star formation. Here one again encounters the $\\approx 60\\Ms$ Eddington limit. Stellar formation models lead to a mass limit near $40-100\\Ms$ imposed by feedback on a spherical accretion envelope \\citep{kah74,wolfc87}. Some observations suggest that proto-stars may be accreting material in disks rather than spheres (\\eg, \\citealt{chi2004}) in which case it may be possible to overcome the radiation pressure at the equator of the proto-star. Studies on the formation of massive stars through disk-accretion with high accretion rates, thereby allowing the radiation to escape preferentially along the poles (\\eg, \\citealt{jija96}) indeed allow formation of stars with larger masses.  The feedback-induced mass limit can also be avoided if massive stars can form through mergers \\citep{bonn98a,zy2007}. In this scenario, massive stars form via coalescence of intermediate-mass proto-stars in the cores of dense stellar clusters that have undergone core-contraction due to rapid accretion of gas with low specific angular momentum. The high central density required for the mergers ($10^8\\Ms$ pc$^{-3}$) is still difficult to achieve but it should be noted that an observable young cluster is necessarily disposed of a substantial fraction of its natal cloud and is thus likely to be always observed in an expanding and hence diluted phase. Recently-set limits on the radii of embedded clusters indeed suggest them to form very compact \\citep{mrk2012}. In this context, it is worthwhile to note that in the computations discussed in the following sections we do get binary coalescence events even in moderately dense stellar clusters. This happens due to the presence of a substantial number of binaries in the cluster core which have much larger encounter cross sections than single stars due to their bigger geometrical sizes.  While the existence of an upper limit of the stellar mass is rather obvious from theory since several decades, such an upper limit has been established from observations only recently. There have been some earlier indications of the presence of a cut-off near $\\imfmax\\approx 150\\Ms$ from observations of young massive clusters, particularly in R136 \\citep{mhunt98,massi2003} in the LMC and in Arches cluster \\citep{fig2003} close to our Galactic center, but these results were not sufficiently convincing in the sense that no statistical significance were attached to them. The observed upper limit was considered to be a limitation due to sampling rather than a true limit \\citep{massi2003}. \\citet{elm2000} also noted that random sampling from an unlimited IMF for all star-forming regions in the Milky Way would lead to the prediction of stars with masses $\\gtrsim 1000\\Ms$, unless there is a sharp down-turn in the IMF beyond several $100\\Ms$.  \\begin{figure*} \\centering \\includegraphics[height=6.0cm,angle=0]{Msc3.eps}\\\\ \\vspace{-0.5cm} \\includegraphics[height=6.0cm,angle=0]{Msc4.eps} \\vspace{-0.5cm} \\caption{Formation of SC stars in the computed ``Model 3'' (upper panel) and ``Model 4'' (lower panel) of \\citet{bko2011}. The panels show the masses of the SC stars at the times of their first appearance and their subsequent mass depletion due to stellar wind as obtained from within NBODY6. For ``Model 3'' (upper panel), the SC stars represented by the grey lines, that appear from the very beginning, are the ``spurious'' ones in this example in the sense explained in Sec.~\\ref{xtracomp}. Similar description applies for the lower panel.} \\label{fig:sc3} \\end{figure*}  \\citet{wk2004} for the first time gave a critical look at the question of the existence of a physical upper-limit of $\\approx 150\\Ms$ in the IMF, as observed in R136 (the central massive star cluster in the 30 Doradus region of the LMC). Assuming R136 has a mass in the range $5\\times 10^4 < M_{cl} < 2.5\\times 10^5\\Ms$ \\citep{sel99}, they found that for a canonical IMF with $\\imfmax=\\infty$, $10 < N(>150\\Ms) < 40$ stars are missing in R136 that have masses $>150\\Ms$. The probability that no stars are observed among the 10 expected ones, assuming $\\imfmax=\\infty$, is $p=4.5\\times10^{-5}$, \\ie, the observed massive stellar content of R136 implies a physical stellar mass limit at $\\imfmax\\approx150\\Ms$. Similarly, \\citet{fig2005} found a dearth of $N(>150\\Ms)=33$ stars in the Arches cluster, where the shallowing of its stellar mass function due to its rapid tidal dissolution has been incorporated. The corresponding number of missing stars for a canonical IMF is $N(>150\\Ms)=18$ which gives a non-detection probability as small as $p=10^{-8}$. Thus the Arches cluster also has $\\imfmax\\approx150\\Ms$ with a very high significance. Finally, \\citet{oey2005} studied 9 clusters and associations in the Milky Way, LMC and SMC to investigate the expected masses of the most massive stars in these for different upper mass limits (120, 150, 200, 1000 and $10000\\Ms$). They concluded that the observed number of massive stars supports the existence of a general physical upper mass cutoff within the range $120\\Ms < \\imfmax < 200\\Ms$ with a high significance.  The massive stellar population of young massive star clusters therefore indicates a physical upper mass limit near $\\imfmax\\approx150\\Ms$ which is the canonical upper limit of the stellar IMF. One can regard this mass as the limit imposed by the process of star formation under physical conditions achievable in a star-forming core. The origin of this very limit is under investigation and the value should currently be taken as empirical only. As coined by \\citet{pk2011}, a stellar population containing stars with their zero-age-main-sequence (ZAMS) masses up to $150\\Ms$ is ``saturated''.  In this paper, our concern is again related to the massive stellar population in R136. The particular aspect that we focus on is related to a recent study by \\citet{crw2010}. These authors re-analysed the massive stellar population of R136 in unprecedented detail using Hubble Space Telescope and Very Large Telescope spectroscopy and high spatial resolution near-IR photometry to find 4 stars, within the central $1\\times1$ pc of R136, with masses $165 - 320\\Ms$, \\ie, substantially above the canonical limit One can call such a stellar population, containing single stars with masses substantially exceeding $150\\Ms$, as ``super-saturated'' and the stars surpassing the canonical upper limit can be called ``super-canonical'' (\\citealt{pk2011}; hereafter SC). Although the stellar population of R136 has been studied by earlier authors, the crucial turn in \\citet{crw2010} is due to these authors' consideration that the observed SC stars in R136, in spite of exhibiting WN-type spectra and possessing strong winds, are actually young, main-sequence stars rather than classical Wolf-Rayet (WR) stars. To take this into account, \\citet{crw2010} incorporate detailed stellar evolutionary models of very massive main-sequence stars that include state-of-the-art treatment of stellar winds which lead to the inference of the SC masses.  More recently however, \\citet{bko2011} (hereafter Paper I) showed that single stars with masses $m_s>150\\Ms$ can indeed form in R136 through mergers of the members in massive binaries. These authors computed the evolution of model clusters that resemble R136 using the state-of-the-art direct N-body integration code ``NBODY6'' \\citep{ar2003}. In these computationally challenging models, all the O-stars were taken to be in tight binaries as observations indicate (\\citealt{sev2010}; see Sec.~\\ref{compute}) but the ZAMS mass of a single component never exceeds $\\imfmax=150\\Ms$. They found that in course of their dynamical evolution, the massive binaries in the model clusters can often merge due to hardening and/or eccentricity enhancement due to dynamical encounters to produce single stars of $m_s>150\\Ms$. This implies that the presence of SC stars in R136 does not illustrate a violation of the canonical upper limit.  Of course, the value of the upper stellar mass limit is not necessarily exactly $=150\\Ms$ \\citep{oey2005}. However, for definiteness, we take the following \\emph{hypothesis}: there exists a fundamental upper mass limit of $\\imfmax=150\\Ms$ or the canonical limit of ZAMS stars formed in clusters and any observed more massive (``super-canonical'') star is created from dynamically-induced stellar mergers within the dense young cluster. In this study, we aim to test the above \\emph{hypothesis} or answer the following question: \\emph{Can the observed number of super-canonical stars in R136 be explained by the naturally occurring stellar-dynamical and stellar-evolution processes in R136?} To that end, we utilize the 4 direct N-body computed massive cluster models of Paper I mimicking R136 (Sec.~\\ref{initcond} \\& \\ref{evolmethod}) and additionally perform 5 more similar computations (Sec.~\\ref{xtracomp}) to trace the formation of SC stars through dynamical means. These computed models initiate with properties that are consistent with the observed properties of young clusters, in particular, primordial mass segregation \\citep{lit2003,chen2007} and massive stellar binaries with component mass ratio close to unity \\citep{sev2010}. Furthermore, we utilize detailed rotating stellar evolution models in the SC mass range to estimate the mass evolution of the dynamically formed SC stars and the resulting lifetimes in their SC phases (Sec.~\\ref{scwind}). We conclude the paper by summarizing our results and pointing out the limitations of the present study (Sec.~\\ref{discuss}). ",
        "conclusions": "\\label{discuss}  \\begin{table*} \\centering \\begin{minipage}{138mm} \\centering {\\large \\caption{\\large Table summarizing the formation of SC stars in our computed models. The descriptions of the columns are as follows: Col.~(1): model ID, Col.~(2): time $T_0$ at which a SC star first appeared in the model, Col.~(3): birth mass $M_0$ of the first-comer SC star, Col.~(4): birth mass $M_{\\rm max}$ of the most massive SC star formed in the model, Col.~(5): formation time, $T_{\\rm max}$, of the most massive SC star, Col.~(6): maximum number of SC stars, $\\mathcal{N}_{\\rm max,in}$, that remained simultaneously bound to the cluster over a $\\tsc=1.5$ Myr period within $<3$ Myr cluster age, Col.~(7): total number of SC stars, $\\mathcal{N}_{\\rm tot}$, formed over the computation (all of them do not necessarily appear simultaneously or remain bound to the cluster). For Models 1-4 (from \\citealt{bko2011}) the ``spurious'' SC stars (see Sec.~\\ref{xtracomp}) are excluded and only those SC stars that are formed later in these computations are considered.} \\label{tab1} \\begin{tabular}{@{}ccccccc} \\hline Model ID & $T_0$ (Myr) & $M_0$ ($\\Ms$) & $M_{\\rm max}$ ($\\Ms$) & $T_{\\rm max}$ (Myr) & $\\mathcal{N}_{\\rm max,in}$ & $\\mathcal{N}_{\\rm tot}$\\\\ \\hline 1 & 2.6 & 193.9 & 193.9 & 2.6 & 1 & 2\\\\ 2 & 2.0 (3.0)$^a$ & 155.2 (181.4)$^a$ & 181.4 & 3.0 & 1 & 2\\\\ 3 & 0.7 & 236.8 & 246.0 & 1.5 & 4 & 5\\\\ 4 & 1.2 & 172.5 & 206.2 & 2.6 & 1 & 2\\\\ C2 & 1.4 & 220.6 & 220.6 & 1.4 & 1 & 2\\\\ C5 & 1.3 & 224.0 &  224.0 & 1.3 & 3 & 3\\\\ C10$^b$ & 1.2 (2.1)$^a$ & 152.4 (162.5)$^a$ & 225.9 & 2.2 & 2 & 4\\\\ \\hline \\end{tabular} \\footnotetext[1]{The first-comer SC star's mass is too close to $150\\Ms$. The time and the mass corresponding to the next SC appearance is also shown in the parentheses.} \\footnotetext[2]{The remaining two of the additionally computed models (with initially only circular binaries) produce SC stars only with masses marginally above $150\\Ms$.} } \\end{minipage} \\end{table*}  Table~\\ref{tab1} summarizes the SC star formation in the computed models of Paper I and in the new computations. Our general conclusions from the computations are as follows: \\begin{itemize}  \\item Formation of SC stars due to dynamically induced mergers of massive binaries are common in a R136-type cluster. In most of the computations we find multiple SC stars formed within $<3$ Myr which is appropriate for R136 (\\citealt{and2009}; see Figs.~\\ref{fig:sc3} \\& \\ref{fig:circbin}). Most SC stars are formed single while a few of them initially form in wide binaries.  \\item The SC stars typically begin to appear from $T_0\\approx 1$ Myr cluster age or even earlier (\\cf Table~\\ref{tab1}). They tend to form with equal likeliness over a cluster age of 1 - 3 Myr.  \\item The most massive SC star formed in a given computed model is typically close to $M_{\\rm max} \\approx 200\\Ms$ and the most massive one formed is $\\approx 250\\Ms$ (``Model 3''; \\cf Table~\\ref{tab1}). The cluster age $T_{\\rm max}$ corresponding to the formation of the most massive SC star is typically well within $<3$ Myr (\\cf Table~\\ref{tab1}).  \\item Multiple SC stars are occasionally found to exist simultaneously near the cluster's center which are bound to the cluster, over a representative SC phase lifetime of $\\tsc=1.5$ Myr, within $<3$ Myr (\\cf Table~\\ref{tab1}) cluster age. This, along with the second point above, implies that it is quite plausible that R136 harbours multiple SC stars at the present day.  \\item Some SC stars are formed with runaway velocities and escape (\\cf Sec.~\\ref{scform}; Fig.~\\ref{fig:scrad}).  \\end{itemize}  These conclusions conform with the observations by \\citet{crw2010} who found 4 SC stars in R136 in the initial mass range of $165-320\\Ms$. In our computations, SC stars of up to $\\approx 250\\Ms$ are formed which is consistent with the observations taking into account the large uncertainties in the stellar evolution models in this mass range. Also, only one of the models (``Model 3'') has up to 4 SC stars simultaneously bound to the cluster (over a 1.5 Myr period) although other models also host multiple SC stars close to the cluster's center. Some models however contain only one bound SC star within 3 Myr (\\cf Table~\\ref{tab1}).  At this point, it is worthwhile to note that our above conclusions depend somewhat on the age of the R136 cluster which is, of course, not fully settled yet. While the bulk of R136 is $\\approx 3$ Myr old \\citep{and2009} the high-mass stellar population might be younger \\citep{mhunt98,dkot98}. In our computations, the SC stars typically begin to appear from $\\approx1$ Myr cluster age and several of them appear by $\\approx 2.5$ Myr. This time-frame is quite consistent with the possible $\\lesssim 2$ Myr age of R136's massive stellar population given the wide uncertainties in the age of such a young cluster particularly for the massive end of the IMF. There has so far been no direct estimate of the age of the massive end of the stellar population of R136 and the above age-limit is based only on comparisons with the stellar populations of Car OB1 and NGC 3603 \\citep{crw2010}. Given that the stellar IMF of R136 continues to maintain the canonical law \\citep{pk2001} from the low to the high mass range ($1.1\\Ms-120\\Ms$; \\citealt{mhunt98,and2009}), we find it more natural to consider that the whole stellar population of the R136 cluster has formed in a single starburst event whithout a significant age spread. The issue can, of course, be more resolved with better understanding of the evolution of very massive stars and their winds. Finally, a recent study by \\citet{chi2012} shows that the O-star binary distribution can, in fact, be even harder (\\ie, more bound or tighter) than that considered in our computations (see Sec.~\\ref{initcond}). This would lead to the formation of SC stars even earlier which constitutes an important future study.  It is important to remember that the formation of SC stars at early ages in the above computations is facilitated by the adopted initial complete mass segregation. This condition subjects the massive binaries to strong dynamical encounters from the beginning of the cluster evolution. Note however that primordial mass segregation is inferred to have been true for several Galactic globular clusters \\citep{bg2008,mrk2010,strd2011} and open clusters \\citep{lit2003,chen2007}. Notably, for a completely unsegregated model, it would take $\\approx 10$ Myr for the massive binaries to segregate to the cluster's center which makes the early formation of SC stars unlikely. However, the mass segregation timescale shortens substantially with increasing compactness of the cluster. In this context, an important outlook would be to study the formation of SC stars with varying initial compactness and degree of primordial mass segregation.  The drawbacks of the computed R136 models, as discussed in Paper I (see Sec.~4 of Paper I), naturally carry over to the present analyses. These limitations however do not crucially affect SC formation. In particular, the truncation of the binary distribution at $m_s\\approx 5\\Ms$ does not influence the mergers of the most massive binaries, the latter being much more centrally concentrated due to mass segregation. Also, the exact period distribution of the O-star binaries is not instrumental for the formation of the SC stars as long as O-stars are largely found in tight binaries \\citep{sev2010}. Notably, our adopted range of the orbital period for the O-star binaries is similar to that reported by \\citet{sev2010} (see Paper I). The formation of SC stars is generally similar for models with initially thermally distributed eccentricities (Paper I) and for the newer ones with initially circular binaries (except for the spurious SC stars in the former models). This implies that the initial eccentricity distribution of the massive binaries does not crucially influence the formation of SC stars.  The most consistent and accurate way to address the above discussed concerns regarding the effects of a stronger wind mass loss due to the newer stellar evolutionary models is to incorporate the latter in a direct N-body code which is well beyond the scope of this paper. No such direct N-body code exists at the present time. Furthermore, there are substantial uncertainties in the physics of mergers of massive stars at the present time. Therefore, the conclusions in this work are the best one can draw given the current technical limitations. Perhaps it would be possible to do such work in future with a distributively computing, highly modular N-body calculation framework such as ``MUSE'' \\citep{pz2009}. In spite of all these uncertainties it is worthwhile to note that if one focuses to the particular case of R136, then (a) the existence of SC stars is \\emph{strongly supported by observations} and (b) considering the widely used $\\approx 3$ Myr age of R136, the SC stars must have formed later than the birth of R136. Our computations are consistent with the formation of SC stars through dynamically induced mergers of massive binaries. While there are technologically-limited drawbacks in our present analysis, as elaborated in Sec.~\\ref{scobs}, none of these shortcomings dictate to the non-observation of SC stars in R136. Therefore, we can say that we have justified the \\emph{hypothesis} of Sec.~\\ref{intro} through our detailed N-body modelling of R136 with initial conditions chosen in accordance with observations of young star clusters. In other words, from our realistically modelled computations of R136-like star clusters, and from our analyses as presented above, it can now be more definitely concluded that the observed super-saturated stellar population in R136 does not imply a violation of the canonical stellar upper mass limit near $\\imfmax=150\\Ms$."
    },
    "1208/1208.4402_arXiv.txt": {
        "abstract": "We present the first 1.3\\,mm (230\\,GHz) very long baseline interferometry model image of an AGN jet using closure phase techniques with a four-element array. The model image of the quasar 1924-292 was obtained with four telescopes at three observatories: the James Clerk Maxwell Telescope (JCMT) on Mauna Kea in Hawaii, the Arizona Radio Observatory's Submillimeter Telescope (SMT) in Arizona, and two telescopes of the Combined Array for Research in Millimeterwave Astronomy (CARMA) in California in April 2009. With the greatly improved resolution compared with previous observations and robust closure phase measurement, the inner jet structure of 1924-292 was spatially resolved. The inner jet extends to the northwest  along a position angle of $-53^\\circ$ at a distance of 0.38\\,mas from the tentatively identified core, in agreement with the inner jet structure inferred from lower frequencies, and making a position angle difference of $\\sim 80^{\\circ}$ with respect to the cm-jet. The size of the compact core is 0.15\\,pc with a brightness temperature of $1.2\\times10^{11}$\\,K.  Compared with those measured at lower frequencies, the low brightness temperature may argue in favor of the decelerating jet model or particle-cascade models. The successful measurement of closure phase paves the way for imaging and time resolving Sgr A* and nearby AGN with the Event Horizon Telescope. ",
        "introduction": "The quasar 1924-292 (PKS 1921-293, OV-236) is one of the brightest and most compact flat-spectrum radio sources in the sky. It has been classified as an optically violent variable \\citep{1981Natur.289..384W,1988AJ.....96.1215P} and highly polarized quasar \\citep{1990ApJ...360..396W}. As a radio-loud blazar, it shows strong variability from radio to X-ray. This source is also included in the Fermi-LAT 1-year Point Source Catalog \\citep{2010ApJS..188..405A}. At its redshift of z=0.352 \\citep{1981Natur.289..384W}, an angular resolution of 1\\,mas corresponds to 4.93 pc ($H_{\\rm 0}= 71 \\rm km s^{-1}  Mpc^{-1}$, $\\Omega_{\\rm M} = 0.27$ and $\\Omega_\\Lambda = 0.73$).  1924-292 is completely unresolved with Very Large Array (VLA) observations made at 6\\,cm and 20\\,cm \\citep{1985AJ.....90..846D,1982AJ.....87..859P}. Very long baseline interferometry (VLBI) observations made at cm wavelengths show a typical core-jet structure with the jet extending about 10 mas to northeast along a position angle (P.A.) of approximately $25^{\\circ}$--$30^{\\circ}$\\citep[e.g.,][]{1998AJ....115.1295K,1989AJ.....98....1P,1997AJ....114.1999S,1998ApJ...497..594T}. With increased resolution, VLBI observations at 7 and 3.5~mm and VLBI space observatory program observations showed that the inner jet curves sharply and is oriented toward the northwest (extends up to about 1 mas) with a time-varying position angle \\citep{2008AJ....136..159L,1999PASJ...51..513S,2002aprm.conf..401S}. \\citet{2002aprm.conf..401S} reported that superluminal motion (about 3\\,c) has been detected. The innermost ($<$ 1\\,pc) region, however, was characterized by two equally compact components whose relative positions were unchanged over about 6.5 years covered by their observations. At cm wavelengths, this source also has one of the highest brightness temperatures of $\\gtrsim 3 \\times 10^{12}$ K measured in the source rest frame, in excess of the inverse Compton limit for synchrotron radiation \\citep{1989ApJ...336.1105L,1996AJ....111.2174M,1999PASJ...51..513S,1998ApJ...497..594T}.  VLBI observations at short millimeter wavelengths ($\\lambda \\leq$1.3\\,mm, $\\nu \\geq 230$\\,GHz) have traditionally been challenging due to the limited sensitivity of  the instruments and  atmospheric phase fluctuations. With the application of new technical developments (such as phased-array processors, wide-bandwidth digital backends, and high-data-rate recorders) and the appearance of new suitable antennas, recent observations have established the technical feasibility of VLBI at short millimeter wavelengths and have opened a new window to directly study and image black holes with the Event Horizon Telescope (EHT)~\\citep{Doeleman2008,2011ApJ...727L..36F}.  Closure phase, which is the sum of the interferometric phase around a triplet of antennas, is largely immune to atmospheric and instrumental complex gain variations~\\citep{1974ApJ...193..293R,1984ARA&A..22...97P}. It is closely related to the asymmetry of the emission and thus is a robust observable for understanding source structures with high resolution. In the low signal-to-noise (SNR) regime, however, it has been inherently difficult to obtain closure phase with high-frequency VLBI at 1.3\\,mm. With the deployment of new VLBI systems and enhancement of software capabilities, closure phases have now been robustly measured by the EHT on the quasar 1924-292, allowing us to model the compact structure on submilliarcsecond scales. ",
        "conclusions": "\\label{discussion} \\begin{figure}[ht] \\begin{center} \\includegraphics[width=0.45\\textwidth,clip]{fg4.eps} \\caption{Model image of 1924-292. Contours are drawn at 1, 2, 4, ..., 64\\,\\% of the peak brightness. The two dashed curves indicate (schematically) how the inner jet is bent toward the cm jet.} \\label{Fig:map} \\end{center} \\end{figure}  There is broad consistency among the three models we have considered: They all require two compact components, $\\sim$ 20--30\\,$\\mu$as (Ma1/Mb1/Mc0) and $\\sim$ 50--80\\,$\\mu$as in size, separated by  $\\sim$370\\,$\\mu$as at a P.A. of $-54^{\\circ}$. These two components are the most reliably measured features by our data, and the accurately measured closure phases remove the $180^{\\circ}$ of degeneracy in the jet position angle.  Of the three models, model Mc gives the best fit based on the $\\chi^2$-fitting results; however, we cannot completely rule out the first class of models. From Figures~\\ref{Fig:radplt} (right) and \\ref{Fig:clplt}, it is obvious that some extra data samples on the existing baselines would easily discriminate between the three models even with the present-day sensitivity. Nevertheless, with the most compact component being the core and the sizes of the three components increasing with distance from the core, model Mc is more consistent with the expectation of an expanding jet. Model C also accounts for all the flux in specific components and therefore is preferred. We show its model image in Figure~\\ref{Fig:map}. The three-component jet model may be approximating a continuous expanding jet structure, perhaps with nonuniform brightness. Sensitivity and uv coverage limitations currently prevent us from being able to model such a structure, but future observations will be able to directly image the jet structure.  The fraction of the core flux density $\\frac{S_{\\rm core}}{S_{\\rm total}}$, which indicates the compactness of a source,  is $\\sim$ 20\\,\\% at 1.3\\,mm from our observations.  The core, however, is very compact and strong ($\\gtrsim$2\\,Jy), making it suitable as a fringe finder for 1.3 mm VLBI observations. It can be seen from Table~\\ref{Table:model} and Figure~\\ref{Fig:map} that the inner jet structure determined by 1.3 mm VLBI agrees well with the reported inner jet orientation at lower frequencies~\\citep[e.g.,][]{2002aprm.conf..401S}. The three-component model at 3.5~mm by \\citet{2008AJ....136..159L} does not fit the jet structure seen at 1.3~mm, and the P.A. between their two components away from the core is about $70^{\\circ}$. We note, however, that the 3.5~mm observations were made at a different epoch, and the significance of these components is unclear given the limited dynamic range of the 3.5~mm map.  The component (Ma2/Mb2/Mc2) may serve as a link to the extended cm-jet toward a P.A. of $\\sim 30^{\\circ}$ and can be interpreted as the result of a local bend in the jet toward the observer (Figure~\\ref{Fig:map}). Interestingly, the innermost component of about 160\\,$\\mu$as (0.8 pc) from the assumed core in model Mc (component Mc1) extends to the northwest along a P.A. of $-70^{\\circ}$. Its size lies well between the sizes of component Mc0 and Mc2, consistent with the expectation of a fanning-out of the jet. With this component, the innermost jet seems to bend more sharply with respect to the cm emission than previously known with VLBI at lower frequencies, indicating that the jet at these sub-pc scales is extremely curved, reminiscent of a helical jet structure. Furthermore, if this component is associated with the recent mm flare started in 2008\\footnote{http://sma1.sma.hawaii.edu/callist/callist.html?plot=1924-292}, and assuming a time lag of 0.1 yr between component ejection and onset of a mm-flare~\\citep{2003ASPC..299..249K}, we obtain an estimate of the jet speed of 0.2 mas/yr (4 c), slightly faster than, but still consistent with, the reported jet speed of 3\\,c~\\citep{2002aprm.conf..401S}. The compact double within the inner 1 pc regions reported by \\citet{2000aprs.conf..155S, 2002aprm.conf..401S} may be associated with either Mc0 and Mc1 or Mc0 and Mc2. However, cross identification is difficult because of the long gap in time between these observations.  The radio core is very compact, and the measured angular size of 31\\,$\\mu$as translates into a linear size of 0.15\\,pc, or $4.6\\times10^{17}$\\,cm. For this component, we obtained a source-frame brightness temperature of $1.2\\times10^{11}$K, less than the inverse Compton limit \\citep{1969ApJ...155L..71K} and the equipartition limit \\citep{1994ApJ...426...51R}. Therefore, the measured core brightness temperature alone does not readily imply relativistic beaming, although the reported superluminal motion is suggestive of a beaming effect. Indeed, a very large Doppler factor (> 80) was reported in this source \\citep{1999PASJ...51..537F}.  The core brightness temperature at 1.3\\,mm is well below those measured at cm wavelengths \\citep[$\\gtrsim 3 \\times 10^{12}$ K,][and references therein]{1999PASJ...51..513S}. This is even true when we consider those derived brightness temperatures as lower limits because the core is unresolved at cm wavelengths. On the one hand, this may suggest that relativistic beaming does not play a significant role in the core region at 1.3\\,mm. One possibility is that the jet bends such that the Doppler beaming effect is not maximized at the position of the core at 1.3\\,mm, but somewhere downstream of the jet.  Another possibility is that the jet undergoes parsec-scale acceleration, so that the jet is gradually beamed at some distance from the central engine, and the inner jet traced at 1.3\\,mm is still accelerating. On the other hand, our observations are at frequencies above the turnover frequency of the synchrotron jet emission for this source. The self-similarity of the parsec-scale jet is believed to break down in the regions probed at these high frequencies~\\citep{2009arXiv0909.2576M}. Compared with previous low-frequency observations, we probed ``deeper'' into the core region, where the brightness temperature can be intrinsically lower. This supports the decelerating jet model or particle-cascade models, as discussed by \\citet{1995PNAS...9211439M}, which predict a lower brightness temperature for the inner jet close to its origin."
    },
    "1208/1208.1157_arXiv.txt": {
        "abstract": "We present Swarm-NG, a C++ library for the efficient direct integration of many \\nbody systems using a Graphics Processing Unit (GPU), such as \\nv's Tesla T10 and M2070 GPUs. While previous studies have demonstrated the benefit of GPUs for \\nbody simulations with thousands to millions of bodies, Swarm-NG focuses on \\emph{many few-body} systems, e.g., thousands of systems with 3\\ldots 15 bodies each, as  is typical for the study of planetary systems. Swarm-NG parallelizes the simulation, including both the numerical integration of the equations of motion and the evaluation of forces using \\nv's ``Compute Unified Device Architecture'' (CUDA) on the GPU. \\jorg{ Swarm-NG includes optimized implementations of 4th order time-symmetrized Hermite integration and mixed variable symplectic integration, as well as several sample codes for other algorithms to illustrate how non-CUDA-savvy users may themselves introduce customized integrators into the Swarm-NG framework.} To optimize performance, we analyze the effect of GPU-specific parameters on performance under double precision. For an ensemble of  131072 planetary systems, each containing 3 bodies, the \\nv\\ Tesla M2070 GPU outperforms a 6-core Intel Xeon X5675 CPU by a factor of $\\sim 2.75$. Thus, we conclude that modern GPUs offer an attractive alternative to a cluster of CPUs for the integration of an ensemble of many few-body systems.  Applications of Swarm-NG include studying the late stages of planet formation, testing the stability of planetary systems and evaluating the goodness-of-fit between many planetary system models and observations of extrasolar planet host stars (e.g., radial velocity, astrometry, transit timing). While Swarm-NG focuses on the parallel integration of many planetary systems, the underlying integrators could be applied to a wide variety of problems that require repeatedly integrating a set of ordinary differential equations many times using different initial conditions and/or parameter values. ",
        "introduction": "\\subsection{Background} An \\nbody simulation numerically approximates the evolution of a system of bodies in which each body continuously interacts with every other body, a fundamental component of many physical and chemical systems. \\nbody simulations are ubiquitous in astrophysics and planetary science.  Example applications include investigating the trajectories of spacecraft, the formation and orbital evolution of the solar system and other planetary systems, the delivery of water to Earth via collisions with asteroids and/or comets, the evolution of star clusters, the formation of galaxies and even the evolution of the entire universe.  Given the widespread use of \\nbody simulations, astronomers have developed a variety of algorithms and computer programs for performing \\nbody integrations. At its core, an \\nbody simulation requires solving a set of $3\\times\\nbod$~ second-order ordinary differential equations (ODEs). For $\\nbod\\ge3$, most sets of initial conditions result in chaotic evolution and are best studied numerically. The computational requirements of \\nbody simulations can be significant, either due to long timescales (e.g., billions of years), a large number of bodies (e.g., $\\sim10^{5\\ldots 7}$ for star cluster, $>10^{9}$ for galaxy or universe), and/or the need to consider a large number of systems with slightly different initial conditions (e.g., $\\sim10^{6\\ldots 9}$ model evaluations for Bayesian analysis of exoplanet observations).  Several previous studies have demonstrated that the combination of modern Graphical Processing Unit (GPU) hardware and the CUDA (Compute Unified Device Architecture) programming environment can greatly accelerate a gravitational \\nbody problem for large $\\nbod$ (e.g., star clusters) \\cite{ 2012JCoPh.231.2825B, Belleman2008, Capuzzo-Dolcetta2011, Gaburov2009, Konstantinidis2010, Hubert:2007:Gem3, Nitadori2012, Zwart2007}. Here we focus on the problem of integrating an ensemble of \\emph{many few-body} systems, e.g., thousands to millions of systems with 3 to tens of bodies each, as is needed for the study of planetary systems.  Given the large number of integrations of few-body systems, GPUs can dramatically reduce the total time required to obtain scientific results for many real-world applications (\\S\\ref{Sec:Applications}).            In this paper, we present the Swarm-NG library for parallel integration of \\nbody systems. In \\S\\ref{Sec:Problem}, we describe the physical setup and the numerical methods. In \\S\\ref{Sec:Cuda}, we describe key aspects of GPU computing with CUDA. In \\S\\ref{Sec:Implementation}, we discuss the details of our implementations. In \\S\\ref{Sec:Results}, we present performance benchmarks. In \\S\\ref{Sec:Discussion}, we discuss present and likely applications. ",
        "conclusions": ""
    },
    "1208/1208.3963_arXiv.txt": {
        "abstract": "{We have performed RV monitoring of the components of the binary system HD 106515 over about 11 years using the high resolution spectrograph SARG at TNG. The primary shows long-period radial velocity variations that indicate the presence of a low mass companion whose projected mass is in the planetary regime ($m \\sin i = 9.33~M_{J}$).  The 9.8 years orbit results quite eccentric ($e=0.57$), as typical for massive giant planets. Our results confirm the preliminary announcement of the planet included in Mayor et al.~(2011). The secondary instead does not show significant RV variations. The two components do not differ significantly in chemical composition, as found for other pairs for which one component hosts giant planets. Adaptive optics images obtained with AdOpt@TNG do not reveal additional stellar companions. From the analysis of the relative astrometry of the components of the wide pair we put an upper limit on the mass of the newly detected companion of about $0.25~M_{\\odot}$. State of art or near future instrumentation can provide true mass determination, thanks to the availability of the wide companion HD106515B as reference. Therefore, HD106515Ab will allow deeper insight in the transition region between planets and brown dwarfs.} ",
        "introduction": "\\label{s:intro}  The upper mass limit of planetary objects is currently widely debated in the scientific community. On one hand, an operational definition can be based on the minimum mass for deuterium burning (about $13~M_{J}$) as a dividing line between planets and brown dwarfs \\citep{2001RvMP...73..719B}, as defined by IAU \\citep{2012IAUTB..28..138B}, or to different fixed threshold values \\citep[e.g. $24-25~M_{J}$:][]{2006ApJ...646..505B,2011A&A...532A..79S}. On the other hand, a definition based on formation mechanisms is more difficult to obtain. In fact, observational data are incomplete and the detection techniques in most cases can provide only indirect indication on the formation of a detected substellar objects. Furthermore, only minimum masses are known for most of the objects detected by RV surveys or, in the case of objects detected through direct imaging,  true masses are uncertain because of the large sensitivity of luminosity on age and intrinsic uncertainties of theoretical models especially at young ages \\citep{2003A&A...402..701B}.  In spite of these difficulties there is growing evidence for a significant overlap in mass between objects formed like stars do and objects formed in a protoplanetary disks. Several objects of planetary mass have been detected as free floating objects in star clusters, star forming regions or in the field using imaging \\citep[e.g.][]{2000Sci...290..103Z,2012ApJ...748...74L,2012arXiv1207.1449S} and microlensing \\citep{2011Natur.473..349S} or as very wide companions of stars \\citep[e.g.][]{2005A&A...438L..29C}. Some of these objects might be formed in planetary systems and then pushed at very wide separation or ejected from the system because of dynamical interactions with other planets \\citep[Jumping Jupiters scenario,][]{2002Icar..156..570M}. However, it seems unlikely this is the only mechanism producing free floating objects below deuterium burning mass \\citep{2011ApJ...743..148B}. Instead, the minimum mass for core collapse was found to be of a few $M_{J}$ and is then likely that objects of planetary mass formed star-like outside planetary disks \\citep{2007prpl.conf..459W}. The statistics of low mass brown dwarfs appear to be compatible with Jeans mass fragmentation of an interstellar molecular cloud \\citep{2009A&A...493.1149Z}.  On the other hand, there are objects with masses from $13$ to about $25~M_{J}$ that are found in systems with other lower mass planets, such as HD168443 and HAT-P13. In the latter case the planetary nature of the lower mass companion is confirmed by the occurrence of transits \\citep{2009ApJ...707..446B}. In other cases such as HD~38519 and HD~202206 a debris disk is present in the system beside a massive planet and a lower mass planet \\citep{2010ApJ...717.1123M}. These facts support the formation of these objects in a protoplanetary disks. Planet formation models also predict the presence of very massive planets, up to about $38~M_{J}$, in exceptional cases of long-lived, massive and metal-rich disks \\citep{2009A&A...501.1139M}. The rising mass function below about $20-30~M_{J}$ \\citep{2006ApJ...640.1051G} is another indication that a different formation mechanism starts to be present above deuterium burning mass.  This picture is further complicated by evidences that the statistical properties of planetary mass companions with (projected) masses between  $4$ to $15~M_{J}$ are different from those of lower mass planets \\citep{2007A&A...464..779R}. It is then possible that either a different formation mechanism is in action or that the evolution of massive planets in a protoplanetary disk is different depending on planetary mass.  The statistic of substellar objects in the mass range between $10$ to $30~M_{J}$ is still limited, given the intrinsic rarity of these objects \\citep{2010A&A...509A.103S,2012A&A...538A.113D}. The discovery of additional candidates is then welcome, especially when the true mass of the companion can be determined or significantly constrained through astrometry \\citep{2011A&A...525A..95S,2011A&A...527A.140R}. We present here the confirmation of a high-mass planet candidate orbiting the star HD106515A and clues on its mass from astrometry. The planet was first included in the compilation by \\cite{2011arXiv1109.2497M} but only orbital period, RV semiamplitude, eccentricity, and corresponding minimum masses and semimajor axis are listed, postponing a more detailed analysis to a forthcoming paper. HD 106515A is part of a wide binary system, with the companion HD 106515B at a projected separation of about 250 AU. Both components were observed as part of the RV survey looking for planets around the components of moderately wide binaries performed using SARG at TNG. ",
        "conclusions": "\\label{s:discussion}  The planet candidate around HD106515A, with its minimum mass of about $9.5~M_{J}$, is one of the few with projected masses close to deuterium burning limit. From available data, we did not detect any additional planets in the system. The high eccentricity of its orbit and the solar-like metallicity of its parent star are in agreement with the differences in the statistical properties of planets below and above 4 $~M_{J}$ found by \\cite{2007A&A...464..779R}.  The planet host has a stellar companion of similar mass, then HD106515Ab adds to the growing census of exoplanets in multiple systems \\citep{2007A&A...462..345D,2012A&A...542A..92R}. The binary separation is quite wide and plausible orbits leave dynamically stable zones up to 40-80 AU around the stars. This suggests a limited impact of the companion on the planet properties but the moderate eccentricity might also be linked to Kozai interactions, which are effective even for widely separated companion considering the old age of the system \\citep{2005ApJ...627.1001T}. We also found from the analysis of the relative astrometry tentative indication of an additional object with a period of 70 yr. From our data, we can not confirm the reality of this object, that might be a very low mass star, and infer the component around which it should be orbiting.  The presence of a well separated companion with similar properties allowed us to perform a sensitive differential abundance analysis \\citep{2004A&A...420..683D,2006A&A...454..581D}. The lack of significant metallicity differences between the components extends the previous finding that large alterations of chemical abundances somewhat linked to the presence of planets are not a common event \\citep{2006A&A...454..581D,2011A&A...533A..90D}. There are 8 binary systems with giant planets suitable for the comparison of chemical abundances \\citep[HD 106515 and the 7 listed in Table 8 of][]{2011A&A...533A..90D}. All of them have $\\Delta$ [Fe/H] $<0.05$ dex, which is significantly smaller than the typical difference in metallicity between giant planet hosts and nearby field stars \\citep[$\\Delta \\rm{[Fe/H]} \\sim 0.25$ dex, see e.g.][]{2005ApJ...622.1102F,2004A&A...418..989N}. This supports the primordial origin for the metallicity enhancement of stars with giant planets \\citep{2005ApJ...622.1102F}.  The binarity of HD106515 represents also a suitable opportunity for the true mass determination HD106515Ab, thanks to the reference provided by HD106515B. The expected astrometric amplitude is of about 1.2 mas for the minimum mass and about 10 mas at stellar/substellar boundary.  From available relative astrometry, the orbital motion of the wide pair is clearly detected. From the analysis of residuals from the long term trend at epochs 1959-1980 where several high-quality data are available we put an upper limit to the mass of the companion of about $0.25~M_{\\odot}$. Much better astrometric precision can be obtained by more recent instrumentation. HD106515 is an ideal target for differential astrometry using AO systems \\citep{2009MNRAS.400..406H,2010EAS....42..179R} and new interferometric instruments like PRIMA \\citep{2011EPJWC..1607005Q}. Subsequently, the combination of ground-based radial velocities and Gaia high-precision space-borne astrometric data might prove decisive \\citep[e.g.,][and references therein]{2011EAS....45..273S}.  Therefore, there are very promising perspectives for a true mass determination of the companion of HD106515A in the coming years, then removing the ambiguity due to projection effects. The availability of true masses is relevant for a better understanding of the high-mass tail of the planetary mass function  and the transition between planets and brown dwarfs. The direct detection of the companion is instead more challenging even for next generation planet finders as SPHERE or GPI \\citep{2010ASPC..430..231B}, because of the small projected separations ($<0.2$ arcsec) and faint luminosities implied by the old age of the system, unless the orbit is seen nearly pole-on and its mass is significantly larger than the minimum mass."
    },
    "1208/1208.3364_arXiv.txt": {
        "abstract": "{We discuss three different ways in which stellar feedback may alter the outcome of star cluster formation: triggering or suppressing star formation, and redistributing the stellar population in space. We use detailed Smoothed Particle Hydrodynamics (SPH) simulations of HII regions in turbulent molecular clouds to show that all three of these may happen in the same system, making inferences about the effects of feedback problematic.} ",
        "introduction": "\\label{sec:1} To what degree star formation is self--regulating is much debated in astrophysics. Stellar feedback in the form of HII regions, winds, jets, radiation pressure, non--ionizing radiation and supernova explosions are all potentially able to influence the star formation process in molecular clouds (e.g. \\cite{matzner2002}). These processes may have a positive or negative effect on star formation, but positive effects (in the sense of triggering) have received the most observational attention (e.g. \\cite{koenig2008}, \\cite{puga2009}, \\cite{zavagno2010}). However, disentangling how these various feedback mechanisms influence star formation is fraught with difficulty, since it is necessary to think comparatively and infer how star formation would proceed differently if feedback were absent.\\\\ \\indent From this perspective, there are three ways in which feedback may influence the formation of stars. Its effect may be positive (commonly referred to as `triggered star formation'), in the sense of increasing the star formation rate or efficiency, producing more stars, or leading to the birth of stars which would otherwise not exist (note that these effects are not necessarily equivalent and, in the same system, some may transpire while others do not). Feedback may also be negative and do the opposite of these things, which we will call `suppressed star formation'. Of course, the global influence of feedback on a given system may be different from its local effects -- it is perfectly possible for star formation to be suppressed at some locations and triggered in others. Finally, suppression and triggering may cancel each other out and feedback may result in the production of a statistically indistinguishable population of stars, but distribute them differently in position or velocity space relative to their distribution in the absence of feedback, for example in well--defined shells. This could be termed `redistributed star formation'.\\\\ \\indent In order to infer in a given system which of these processes is at work and what is the overall influence on the end product -- the stellar cluster -- it is essential to have a credible counterfactual model for comparison. This is equally true of observed and simulated systems. A good idea of what the system would have done in the absence of feedback can then be gained, and hence the effects of feedback isolated.\\\\ ",
        "conclusions": ""
    },
    "1208/1208.5343_arXiv.txt": {
        "abstract": "\\noindent The AMI Galactic Plane Survey (AMIGPS) is a large area survey of the outer Galactic plane to provide arcminute resolution images at milli-Jansky sensitivity in the centimetre-wave band. Here we present the first data release of the survey, consisting of 868\\,deg$^2$ of the Galactic plane, covering the area $76^{\\circ} \\lessapprox \\ell \\lessapprox 170^{\\circ}$ between latitudes of $|b| \\lessapprox 5^{\\circ}$, at a central frequency of 15.75\\,GHz (1.9\\,cm).  We describe in detail the drift scan observations which have been used to construct the maps, including the techniques used for observing, mapping and source extraction, and summarise the properties of the finalized datasets.  These observations constitute the most sensitive Galactic plane survey of large extent at centimetre-wave frequencies greater than 1.4\\,GHz. ",
        "introduction": "Large-area radio surveys contribute to our understanding of the Universe in numerous and diverse ways. Discoveries from these surveys have become key ingredients of modern astrophysics: pulsars, radio galaxies and quasars and more (see e.g. \\citealt{1998ASSL..226....3L}).  For studies of our Galaxy, radio surveys are particularly beneficial as the longer wavelength radio emission does not suffer from the same extinction and opacity effects as optical and infra-red surveys and the dense regions of dust and gas which dominate the low-latitude Galactic plane become largely transparent, allowing us to study sources in these regions.  However, the bulk of Galactic radio surveys are at frequencies at or below 1.4\\,GHz and as such are necessarily biased against objects whose spectra rise with frequency, such as dense star-forming regions.  Two examples of the need for higher-frequency, centimetre-wave Galactic surveying are as follows.  The first is the hypercompact {\\sc Hii} ({\\sc HCHii}) region. Thought to indicate the earliest visible stage of massive star formation, these objects are two orders of magnitude more dense than the better known ultracompact ({\\sc UCHii}) region and have steeply rising spectra. {\\sc HCHii} regions were discovered serendipitously in observations of {\\sc UCHii}, having been missed previously in their entirety by Galactic plane surveys concentrated at $\\nu<5$\\,GHz. The turnover frequency between the optically thick and thin regimes for thermal bremsstrahlung is a linear function of emission measure (e.g. \\citealt{1967ApJ...147..471M}) causing such low frequency surveys (e.g. $\\nu\\leq 5$\\,GHz) to preferentially select against dense plasmas ($n_{\\rm e} \\leq 10^{11}$\\,m$^{-3}$). Such plasmas are not limited to {\\sc HCHii} regions but also include a variety of other Galactic objects such as massive stellar winds, ionised jets from young stellar objects (e.g. \\citealt{1995RMxAC...1...67A}) and young planetary nebulae (e.g. \\citealt{2009MNRAS.397.1386B}).  The second is the anomalous microwave emission (AME), now being identified in an increasing number of Galactic objects, that was missed in low frequency Galactic surveys. First identified by CMB experiments \\citep{1997ApJ...486L..23L} as a large scale foreground contaminant this form of emission has since been demonstrated to exist in more compact objects such as dark (e.g. \\citealt{2006ApJ...639..951C}; \\citealt{2009MNRAS.394L..46A}; \\citealt{2010MNRAS.403L..46S}) and molecular clouds (\\citealt{2005ApJ...624L..89W}; \\citealt{2011MNRAS.418.1889T}). Although multiple mechanisms have been proposed to explain AME, dipole emission from rapidly rotating very small dust grains (\\citealt{1998ApJ...494L..19D}, \\citealt{1998ApJ...508..157D}) is generally considered to be most likely. Such spinning dust emission has a peaked SED with a maximum in the frequency range 10 -- 50\\,GHz depending on grain size distributions.  A current lack of surveys in this frequency range means that our knowledge of the overall properties of objects which exhibit emission from spinning dust, objects which are characterized by dense plasmas, and indeed the global distribution of rising-spectrum emission in the Galaxy, is extremely poor. Those surveys which are available, such as the 9C Ryle Telescope survey (15\\,GHz; \\citealt{2003MNRAS.342..915W}), the GPA survey (14.35\\,GHz; \\citealt{2000AJ....119.2801L}) and the AT20G survey (20\\,GHz; \\citealt{2010MNRAS.402.2403M}) have provided us with tantalising insights into the high frequency Galactic plane, but there is a continuing need for higher sensitivity, resolution and sky area coverage at these frequencies.  The interferometric Arcminute Microkelvin Imager (AMI) Galactic Plane Survey (AMIGPS) provides the most sensitive centimetre-wave Galactic plane survey of large extent at $\\nu > 1.4$\\,GHz.  AMIGPS is a drift-scan survey of the northern Galactic plane at $\\approx16$\\,GHz, covering (in the first data release) the region $76^{\\circ} \\lessapprox \\ell \\lessapprox 170^{\\circ}$ and $|b|\\lessapprox 5^{\\circ}$.  The AMI Small Array (SA) has been used for the survey since its relatively large field of view ($\\approx400$\\,arcmin$^{2}$) makes covering large areas feasible, and its short baselines mean that extended objects, very common in the Galaxy, are at least partially observable.  The resolution of the survey is $\\approx3$\\,arcmin and the noise level is $\\approx3$\\,mJy\\,beam$^{-1}$ away from bright sources.  This paper focuses on the techniques employed for observing (Section~\\ref{Observations}), mapping (Section~\\ref{Data reduction and mapping}) and source extraction (Section~\\ref{Source extraction}) in the AMIGPS.  The positional and flux density calibration accuracy of the survey are also tested in Section~\\ref{Calibration accuracy checks}, and in Section~\\ref{Data products} the maps and catalogue are described.  In a following paper, hereafter Paper~II, the first results from the survey, including the follow-up of rising-spectrum objects in order to detect \\textsc{U/HCHii} regions, will be presented. ",
        "conclusions": "The Galactic plane between $b\\approx\\pm5^{\\circ}$ has been surveyed using the interferometric AMI SA at $\\approx$16\\,GHz, to a noise level of $\\approx3$\\,mJy\\,beam$^{-1}$ at $\\approx3$\\,arcmin resolution.  This is the most sensitive and highest resolution Galactic plane survey at cm-wave frequencies above 1.4\\,GHz.  \\begin{enumerate} \\item{868\\,deg$^{2}$ of the Galactic plane have been surveyed and a catalogue of 3503 sources produced.  This is the first data release of the AMIGPS.} \\item{As part of creating the AMIGPS, we have developed an automated pipeline to produce maps from data taken in drift-scan mode, accounting for the presence of bright sources.} \\item{The source extraction techniques developed for the 10C survey have been applied to maps at different resolution and regions of the sky with many extended sources present.} \\item{In testing the flux calibration of the survey by comparing source flux densities derived from the AMIGPS to tracked observations of both extra-galactic and Galactic sources taken with the AMI SA and AMI LA, we find that the AMIGPS flux calibration is accurate to within 5\\%.} \\item{The r.m.s. positional accuracy of the survey, assessed by comparing positions derived from the AMIGPS with well-known source positions from the VLBA calibrator survey and with AMI LA follow-up positions, is 2.6\\,arcsec in RA and 1.7\\,arcsec in $\\delta$.} \\end{enumerate}  In a following paper the first results from the survey will be presented, and in a future data release the survey will be extended to $\\delta\\geq20^{\\circ}$."
    },
    "1208/1208.0523.txt": {
        "abstract": "Cosmic-ray and X-ray heating are derived from the electron energy loss calculations of Dalgarno, Yan and Liu for hydrogen-helium gas mixtures. These authors treated the heating from elastic scattering and collisional de-excitation of rotationally excited hydrogen molecules. Here we consider the heating that can arise from all ionization and excitation processes, with particular emphasis on the reactions of cosmic-ray and X-ray generated ions with the heavy neutral species, which we refer to as chemical heating. In molecular regions, chemical heating dominates and can account for 50\\% of the energy expended in the creation of an ion pair. The heating per ion pair ranges in the limit of negligible electron fraction from $\\sim 4.3 $\\,eV for diffuse atomic gas, to $\\sim 13$\\,eV for the moderately dense regions of molecular clouds and to $\\sim 18 $\\,eV for the very dense regions of protoplanetary disks. An important general conclusion of this study is that cosmic-ray and X-ray heating depends on the physical properties of the medium, i.e., on the molecular and electron fractions, the total density of hydrogen nuclei, and to a lesser extent on the temperature. It is also noted that chemical heating, the dominant process for cosmic-ray and X-ray heating, plays a role in UV irradiated molecular gas. ",
        "introduction": "There has been a significant increase of interest in the cosmic-ray ionization rate in both local and distant neighborhoods. Much of it has been spurred by infrared absorption-line measurements of the $\\hthreep$ ion in diffuse and translucent interstellar clouds (e.g., Indriolo et al.~2007 and references therein). Indriolo \\& McCall~(2012) have extended these measurements to fifty lines of sight dominated by translucent clouds with more than one magnitude of visual extinction, obtaining 21 detections. Their straightforward analysis of the observations yields values of the ionization rate for $\\hm$ in the range $\\sim (2-11) \\times 10^{-16}\\,\\ps$, significantly larger than the range of values that had been in use for several decades\\footnote{In this and other papers, the cosmic-ray ionization rate $\\zeta_{\\rm H}$ per H nucleus is used, which is smaller by a factor of two than the rate $\\zeta_{{\\rm H}_2}$ per $\\hm$ molecule of Indriolo et al.\\,(2007).}. This result has led to the reconsideration of how both electron and proton cosmic-rays propagate in the interstellar medium (e.g., Indriolo et al.~2009, Padovani et al.~2009, Padovani \\& Galli 2011, Rimmer et al.~2012, Everett \\& Zweibel 2011). In the context of the new survey, Indriolo \\& McCall al.~(2012) reviewed the explanation advanced by Indriolo et al.~(2009) and Padovani et al.~(2009) that the intensity of low-energy protons is reduced inside dense molecular clouds. The large difference between diffuse/translucent and dense clouds is maintained in Indriolo \\& McCall~(2012), but their sample shows little evidence for a dependence of ionization rate on column density. The value of the cosmic ray ionization rate for the local interstellar medium remains an open question.  Another topic of interest relating to the interaction of cosmic-rays with the interstellar medium is their role in heating. Cosmic-rays are an efficient (often dominant) source of heating in various environments, from the dense gas in molecular clouds (Goldsmith \\& Langer~1978), both in normal and starbust galaxies (e.g., Suchkov et al.~1993), to photodissociation regions (Shaw et al.~2009), and possibly even in the primordial gas (Jasche et al.~2007). In dense, shielded regions like molecular cloud cores, the balance of cosmic-ray heating and gas cooling determines a temperature gradient decreasing inwards (Galli et al.~2002) that has been accurately traced by interferometric observations of molecular emission in prestellar cores (Crapsi et al.~2007, Pagani et al.~2007) and dark globules (Pineda \\& Bensch~2007). The values for the cosmic-ray heating energy per ion pair available in the literature range over a factor of three. In this paper we will attempt to clarify this situation for both cosmic-ray and X-ray heating.  The interaction of fast electrons slowing down in molecular gas was first analyzed more than three decades ago by Glassgold \\& Langer~(1973, hereafter GL73), Cravens et al.~(1975) and Cravens \\& Dalgarno~(1978).  Although GL73 was the first paper to consider cosmic ray heating in interstellar molecular regions, it suffered from the incompletely known electron cross sections of the early 70s, a crude energy-loss calculation, and cumbersome notation. Cravens et al.~(1975) did a better job on cross sections and energy loss, but restricted themselves to low-energy electrons in order to evaluate the validity of the commonly used continuous slowing down approximation. This limitation was removed by Cravens \\& Dalgarno (1978). There are no glaring discrepancies between  Cravens \\& Dalgarno (1978) and GL73. All of these papers considered pure $\\hm$ regions, i.e., the roles of $\\hp$ and $\\hep$ ions in molecular gas were ignored. A more up to date and complete analysis was carried out by Dalgarno et al.~(1999, hereafter DYL). DYL considered carefully all of the energy loss channels for electron energies from 30~eV to 1~keV in various mixtures of $\\h$, $\\hm$ and $\\he$. They showed how the energy expended to make  an ion pair $W$ is partitioned between elastic and non-elastic processes, but they did not fully treat the heating.  Cosmic-rays and X-rays (or the equivalent, fast electrons) produce ions and excited molecules that can interact with the dominant neutral atomic or molecular gas. The products of these reactions carry away a significant amount of the available energy and heat the gas.  We refer to this as chemical heating, and quantify it by a quantity $Q_{\\rm chem}$ defined per ion pair. Chemical heating by cosmic-rays and X-rays occurs in a wide range of applications from diffuse and dense interstellar clouds to pre-stellar cores, protoplanetary disks, and planetary atmospheres. In this paper we use the results of DYL to inquire how the excited species and ions heat a gas mixture of $\\h$, $\\hm$, and $\\he$.  We will show that $\\sim 1/2$ of the energy of X-rays and cosmic rays can go into heating in dense molecular regions and that chemical heating can be the most important part of the heating. We also note that chemical heating occurs when molecular regions are irradiated by far ultra-violet (FUV) photons.  This paper is organized as follows. In Section~\\ref{phys} we outline the physical basis of our analysis and make quantitative estimates in Sections~\\ref{heating} and~\\ref{chem}. We discuss the most important results in Section~\\ref{results}.  The paper ends with a short summary in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc}  Dalgarno et al.~(1999; DYL) made an extensive study of the energy loss of fast electrons in H, He and $\\hm$, He gas mixtures. The electrons are at the heart of the interaction of cosmic rays and X-rays with interstellar and circumstellar matter. Although DYL analyzed essentially all excitation and ionization processes, they only discussed the heating from elastic collisions and rotational excitation of $\\hm$. Starting from the results in DYL, we have extended the scope of cosmic-ray and X-ray heating to include all of the relevant interactions. One of the main conclusions of this study is that heating by fast-electrons depends on the physical properties of the gas, i.e., on the abundance of electrons, $\\hm$ molecules, and heavy atoms and molecules, and also on the total density of hydrogen nuclei. The electron fraction is important because, once it exceeds a certain level, heating by collisions with ambient electrons becomes important. The electron abundance also determines whether the destruction of the $\\hthreep$ ions proceeds by dissociative recombination or ionic reactions, which affects the quantitative amount of chemical heating. Of course the $\\hm$/H abundance ratio is important because the diversity of the energy levels of $\\hm$ offers many more channels for energy loss than atomic H. The dependence on physical conditions means that X-ray and cosmic-ray heating cannot be specified by a single number. This should be clear from the fact that in neutral atomic regions the heating efficiency is only 12\\%, whereas in neutral molecular regions it can reach 50\\% at very high densities.  A wide range of values for the heating per ion pair can be found in the literature on molecular clouds, from 7~eV (Stahler \\& Palla 2004), to 20~eV (Goldsmith \\& Langer 1978, Goldsmith~2001), but not necessarily for the reasons given here, e.g., in the above discussion of Table 6. Many authors adopt Goldsmith's value while others estimate intermediate values of 10--15~eV based on GL73, e.g., Maloney et al.~(1996), Yusef-Zadeh et al.~(2007), and Krumholz et al.~(2011). As shown in Table 4, 13~eV is a good choice for not too dense molecular gas.  Chemical heating also applies to regions exposed to far UV radiation (Dalgarno and Oppenheimer 1974). It was then considered by Barsuhn \\& Walmsley~(1977) and Clavel et al.~(1978), who studied the chemical and thermal equilibrium in dark clouds exposed to far UV radiation and cosmic rays. Clavel et al. explicitly included the contribution of every reaction in their chemical network to heating. Chemical heating has also been widely used in the study of planetary atmospheres (e.g., Roble et al.\\,1987).  The results of this paper are based on the fact that roughly half of the energy generated by cosmic rays and X-rays comes goes into ionization of the gas and roughy half into its excitation, more specifically according to the way the individual processes are treated by DYL. Roughly half of the gas heating comes from the reactions of various ions with the assumed mainly neutral gas. Potentially an equal amount of heating can arise from dissociation and from rotational and vibrational excitation, but the yield from excitation depends on whether the physical conditions are conducive to the quenching of the excited levels. If quenching is inefficient, the levels decay with the emission of radiation, which can escape or be absorbed by dust. This possibility is even more important for the excitation of the singlet levels of $\\hm$, e.g., the $B$ and $C$ levels, and the excitation of the $n=2$ level of H in atomic regions. We have not attempted to follow the fluorescent radiation to determine how much escapes and how much is absorbed. This is an important issue when considering the broader thermal properties of the gas. The treatment of the fluorescent radiation involves radiation transfer and depends on the properties of the dust, i.e., it is application specific. In contrast, our goal has been to treat one relatively well-defined part of the thermal problem of interstellar and circumstellar molecular gas.  In pursuing this goal, we have also ignored the direct interaction of X-rays and cosmic-rays with the dust, of some interest because of the possibility that they may release fast electrons from the dust. Although this occurs, it is relatively unimportant. First of all, the dust cross section per H nucleus for the MRN distribution for the diffuse ISM is $1.6\\times 10^{-21} \\psqcm$. For molecular clouds and disk atmospheres, it will be even smaller. The main inelastic cross sections for a keV electron with $\\hm$ are typically several times $10^{-17} \\psqcm$, so the probability that an incident X-ray or cosmic ray interacts with dust is roughly 1000 time less than with $\\hm$. And when this rare event does occur, no more than 1\\% goes into in photoelectrons, as shown by Dwek and Smith (1996) for the EUV/Xray bands."
    },
    "1208/1208.1426_arXiv.txt": {
        "abstract": "We report the results of a joint \\chandra-\\hst\\ study of the X-ray binary population in the massive, high-density globular cluster NGC~6388.  NGC~6388 has one of the highest predicted X-ray binary production rate of any Galactic cluster.  We detected a large population of 61 \\chandra\\ sources within the half-mass radius with L$_X > 5 \\times 10^{30}$ \\ergs.  From the X-ray colors, luminosities, (lack of) variability, and spectral fitting, we identify five as likely quiescent low-mass X-ray binaries.  Due to the extremely crowded nature of the core of NGC~6388, finding optical identifications to \\chandra\\ sources is challenging. We have identified four blue, optically variable counterparts to spectrally hard X-ray sources, evidence that these are bright cataclysmic variables (CVs).  One showed variability of 2 magnitudes in V, indicative of a dwarf nova eruption.  One other likely CV is identified by its X-ray spectrum (partial covering with high $N_H$) and strong variability, making five likely CVs identified in this cluster.  The relatively bright optical magnitudes of these sources put them in the same class as CV1 in M15 and the brightest CVs in 47 Tuc. ",
        "introduction": "Studying the relationship between X-ray binaries (XRBs) and globular clusters continues to provide interesting clues into cluster dynamical evolution. These clusters have long been thought to be major producers of XRBs, since dense cluster cores promote the type of dynamical interactions that lead to tight binaries such as XRBs \\citep{Clark75,Pooley03}.  These interactions include exchange interactions where a massive stellar remnant such as a heavy white dwarf or neutron star has a close encounter with an existing binary, replacing one of the other stars \\citep{Hills76}, tidal capture by a heavy remnant of a main-sequence star \\citep{Fabian75}, or collisions of neutron stars with red giants, eventually producing ultracompact binaries with neutron stars accreting from white dwarfs \\citep{Verbunt87}.  Cataclysmic variables (CVs) and low-mass X-ray binary (LMXB) systems are common products of these interactions \\citep{Ivanova06,Ivanova08}.  In globular clusters, a significant number of millisecond pulsars (MSPs) have been identified, such as in 47 Tuc \\citep{Camilo00}.  Such systems are the likely progeny of LXMBs \\citep{Archibald09}. Since mass segregation causes the more massive populations of binary stars and degenerate remnants (for instance, MSPs) to be more centrally concentrated, these objects will be strongly tied to the dynamical evolution of the cluster. Most significantly, this increases the likelihood of further interactions of binaries in the core which can then halt the progress of dynamical relaxation, supporting the core against core collapse \\citep{Fregeau08}.  The known distances, reddenings, and ages of globular clusters allow for an analysis of populations of CVs and LMXBs that is not possible for field objects. Using the Hubble Space Telescope (\\hst) in conjunction with the Chandra X-ray Observatory (\\chandra)  large numbers of X-ray sources have been identified in clusters and matched to optical counterparts. For instance, \\citet{Heinke05} detected 300 X-ray sources down to $L_X \\sim 8\\times10^{29}$ \\ergs~ within the half-mass radius of 47 Tuc, 105 of which could be securely classified as quiescent LMXBs (qLMXBs), CVs, MSPs, or chromospherically active binaries (ABs), mostly by using HST to identify and characterize optical counterparts \\citep{Edmonds03,Edmonds03b}.  Similar identifications were also made of 59 of the 79 X-ray sources detected (with limiting $L_X \\sim 10^{29}$ \\ergs) in NGC~6397, a cluster that has undergone core-collapse, by \\citet{Cohn10} in a joint \\hst-\\chandra\\ study.   Higher numbers of interactions in cluster cores lead to larger populations of X-ray sources \\citep{Pooley03}.  This interaction rate can be simply estimated by the encounter rate, $\\Gamma \\propto \\rho_0^2 r_c^3/v_0$ \\citep{Verbunt87b,Maccarone11}.  This quantity is a measure of how often stellar interactions will occur as a function of the cluster's central density, core radius and central velocity dispersion. The numbers of X-ray sources in a cluster, particularly the qLMXBs, MSPs, and brighter CVs, vary directly with this encounter rate, implying strongly that dynamical interactions are the primary source of these populations in globular clusters \\citep{Pooley03,Heinke03,Pooley06}, while fainter ABs seem to be largely primordial in origin \\citep{Bassa04}.    \\subsection{NGC~6388\\label{ngc6388}}  NGC 6388 is a particularly good candidate for the study of the dynamical origin of XRBs in dense-core globular clusters due to its large central density ($10^{5.34}~L_\\odot{\\rm pc}^{-3}$) and core radius ($7.2'' = 0.35~{\\rm pc}$ at 9.9 pc from the Sun) \\citep{Harris96}\\footnote{\\url{http://physwww.physics.mcmaster.ca/~harris/mwgc.dat} for 2010 revision.}. These give it one of the largest values of $\\Gamma$ of any globular cluster in the Galaxy. With such a large encounter rate, we expect it to host a rich population of LMXBs and CVs, and thus to find a large number of X-ray sources. In \\citet{Lugger87}, the U-band surface brightness profile, which is dominated by horizontal branch stars,  was found to be well fit by a single mass King model, suggesting that the core is likely being supported by a large number of primordial binaries. It is possible that some of these primordial binaries could be detected now as ABs.   NGC 6388 also has several features beyond its large encounter rate that make it a cluster worth investigating. For instance, for a cluster of its metallicity (as high as ${\\rm [Fe/H]} \\approx -0.7$ \\citep{Wallerstein07}, it has a rather atypical extreme blue horizontal branch (EBHB) in addition to a blue horizontal branch (BHB), and a very populated red horizontal branch (RHB).  \\citet{Rich97} investigated whether this could be the result of dynamical interactions in the core enhancing the BHB via stripping of the envelope mass.  However they  found no difference in the radial distribution of the RHB and BHB stars and took this as evidence against the dynamical argument.  Also intriguing are the blue hook stars (hotter and less luminous than EBHB stars), investigated in \\citet{Dalessandro08}.  \\citet{Noyola06} report evidence for a modest central cusp in the surface brightness profile and \\citet{Lanzoni07} see a modest central surface density cusp in the inner arcsec, which are interpreted by those groups to be evidence for a central intermediate-mass black hole (IMBH) of mass $5.7\\times 10^3\\msun$.  \\citet{Cseh10} show that their radio nondetection gives an upper limit to the mass of a central IMBH that is less than half the mass predicted by \\citet{Lanzoni07}.  Recently \\citet{Lutzgendorf11} used stellar kinematics to argue for a central IMBH with a mass of 17 $\\pm9\\times 10^3 \\msun$.  \\citet{Miocchi07} suggests that tidal stripping of red giant envelopes by this putative IMBH may be producing the significant population of EBHB stars.  \\citet{Verbunt01} analyzed two ROSAT PSPC observations of NGC 6388 from 1991 and 1992, identifying an X-ray source at $L_X$(0.5-2.5 keV)=$6\\times10^{33}$ ergs/s just outside the half-mass radius.  We do not detect a source in the smaller of Verbunt's two error circles.  This could indicate a transient was active during the ROSAT observations, and faint in quiescence during the \\chandra\\ observations.  It also could indicate that Verbunt's ROSAT astrometry was incorrect, since the luminosity is close to our measured total $L_X$ for the cluster.   Less sophisticated X-ray analyses of our Chandra observation and a 2003 XMM observation, with the goal of constraining the presence of an intermediate-mass black hole at the cluster center, were performed by \\citet{Cseh10,Nucita08}.  Evidence that NGC 6388 contains at least one transient X-ray binary (in agreement with predictions of its high stellar encounter rate) was provided by IGR J17361-4441, a hard X-ray transient located in NGC 6388. It was detected by INTEGRAL on August 11, 2011, with follow up observations by Swift, RXTE, XMM, ATCA, and \\chandra.  It had a peak output in the 0.5-10 keV energy band of $L_X\\sim5\\times10^{35}$ ergs/s \\citep{Bozzo11,Nucita12}.  A short \\chandra\\ target of opportunity observation provided a precise position, inconsistent with any of the cluster X-ray sources we present below, allowing an upper limit on the transient's quiescent luminosity of $10^{31}$ ergs/s \\citep{Pooley11}. In the last Swift observation of Nov.  4, 2011 (before solar proximity forbade observations), this transient was still active at similar luminosities, but the X-ray emission from the cluster was consistent with the normal quiescent output by Jan. 29, 2012 \\citep{Bozzo12}. ",
        "conclusions": "\\label{discussion}  We find a rich population of X-ray sources in NGC 6388.  Our analysis of the density profile for the X-ray sources shows them to be more centrally concentrated than the brightest stars in the cluster. The inferred characteristic mass of 1.36 $\\msun$ for the X-ray sources suggests that many of them are likely to be compact binary stars with degenerate remnant primaries. Many of these were likely produced in close interactions in the cluster's large, dense core.  We examine some of the characteristics of the different populations of the optical sources. The extensive population of blue stragglers was shown to be more massive than the horizontal branch stars. The distinctive extreme blue horizontal branch stars were shown to be no different in distribution than the other horizontal branch stars, confirming the result that dynamical interactions are not a likely cause of the large color distribution in the cluster's horizontal branch.  We identify five sources as bright CVs based on their X-ray luminosities and colors, together with the detection of optical counterparts for four of them, and the detection of X-ray variability for the remaining one.  These counterparts are consistent with bright CVs in optical variation, color, and luminosity.  In fact, we find at least one of the sources (CX9) to exhibit variation on par with that exhibited by dwarf nova outbursts. That would make it only the 15th such object found in globular clusters to date \\citep[see ][]{Servillat11}. The only CV candidate without an identified optical counterpart, CX4, was identified as a likely CV based on its X-ray properties; an unusually hard spectrum, well-fit by partial covering absorption models, luminosity, and variability. We find five sources that have X-ray colors, luminosities, and spectra consistent with what we expect for qLMXBS.  These sources are not found to vary, which is consistent with expectation for sources without a strong power-law component.  The CVs found here are some of the brightest (in X-rays and optical light) yet known in globular clusters, suggesting relatively high mass transfer rates. These are likely the tip of a much larger CV population.  The limit of our detection of CVs was $V_{555}<23.1$, for which we detected 4 bright CVs as optical counterparts to X-ray sources. We can extrapolate the possible size of the CV population in NGC 6388 by comparison with other clusters with bright CVs and a detected CV population. 47 Tuc was found to have at least 22 CVs by \\citet{Edmonds03} looking to 8 magnitudes fainter than the MSTO. After adjusting for distance and reddening, there are 4 CVs in 47 Tuc with magnitudes similar to those in NGC 6388. We can infer that the population of CVs in NGC 6388 is likely as large as that of 47 Tuc. This is not surprising given the similarities between the two clusters and that they both have some of the highest encounter rates of any cluster (see below).  The large number of sources detected in our data is consistent with predictions based on the encounter rate. We can estimate the encounter rate using the approximation given in \\citet{Verbunt03}, $\\Gamma \\propto \\rho_0^2 r_c^3/v_0$. We use values from \\citet{Harris96} (2010 edition) for $\\rho_0$, $r_c$, and $v_0$ for this approximation. In Table \\ref{t:collision}, we give our calculation for this value for the clusters studied in \\citet{Pooley03} and for NGC 6388 and compare these to the number of detected X-ray sources with $L_X>4\\times10^{30}$ \\ergs~in the 0.5-6 keV range.   We have normalized our values of $\\Gamma$ to the mean of the values given for the non-core collapsed clusters in \\citet{Pooley03}. We find them to be roughly consistent, with major differences likely due to our use of the updated Harris catalog. NGC 6388's approximate $\\Gamma$ is seen to be significantly larger than that of all but one of the other clusters. In the table we also compare the number of detected sources reported by \\citet{Pooley03} with the number detected in more recent studies (as noted in the table comments).  The larger numbers detected in more recent studies may be attributed to improved reduction techniques that allow for a better separation of sources in crowded regions. Ultimately, the only two clusters to show a significant difference are $\\omega$ Cen and 47 Tuc, which are most affected by the changes in their values for $\\Gamma$, a result of using updated values for the cluster parameters.  In Fig. \\ref{f:gamma}, we plot the number of detected sources, using the newest values when possible and correcting for the predicted background contamination, versus $\\Gamma$. We note that detections for NGC 6388 and NGC 6093 (M80) are not complete to L$_X>4\\times10^{30}$ \\ergs~ (see \\S\\ref{chandra}), so their detected source numbers are lower limits (indicated with arrows). We fit the relationship of $\\Gamma$ and N with a power law and find it can be described as $N\\propto\\Gamma^{0.55\\pm 0.09}$. This corresponds with the result of \\citet{Pooley03}, where a similar trend was found with their calculated values of $\\Gamma$.  In light of this trend, the detection of X-ray sources in NGC 6388 is a confirmation of the prediction that its large, dense core produces a large number of close stellar interactions, which in turn accelerate the production of close binary systems."
    },
    "1208/1208.1756_arXiv.txt": {
        "abstract": "{} { We study the Star Formation History (SFH) of 210 galaxies members of 55 Hickson Compact Groups (HCG) and 309 galaxies from the Catalog of Isolated Galaxies (CIG).  The SFH traces the variation of star formation over the lifetime of a galaxy, and yields consequently a snapshot picture of its formation. Comparing the SFHs in these extremes in galaxy density, allows us to determine the main effects of Compact Groups (CG) on the formation of galaxies. } { We fit our spectra using the spectral synthesis code STARLIGHT and obtain the stellar population contents and mean stellar ages of HCG and CIG galaxies in three different morphological classes: early-type galaxies (EtG), early-type spirals (EtS), and late-type spirals (LtS). } { We find that EtG and EtS galaxies in HCG show larger contents of old and intermediate stellar populations as well as an important deficit of the young stellar population, what clearly implies an older average stellar age in early galaxies in HCG. For LtS galaxies we find similar mean values for the stellar content and age in the two samples. However, we note that LtS can be split into two subclasses, namely old and young LtS. In HCG we find a higher fraction of young LtS than in the CIG sample, in addition, most of these galaxies belong to groups where most of the galaxies are also young and actively forming stars. The Specific Star Formation Rate (SSFR) of spiral galaxies in the two samples differ. EtS in HCG show lower values of the SSFR, while LtS peak at higher values when comparing with their counterparts in isolation. We have also measured shorter Star Formation Time Scale (SFTS) in HCG galaxies, indicating that they have less prolonged star formation activity than CIG galaxies. We take these observations as evidence that galaxies in CG evolved more rapidly than galaxies in isolation regardless of their morphology. Our observations are consistent with the hierarchical galaxy formation model, which states that CG are structures that formed recently from primordial small mass density fluctuations. From the systematic difference in SFTS we deduce that the HCG have most probably formed $\\sim 3$~Gyr in the past. The galaxies in the HCG are not in equilibrium state, but merging without gas (i.e., under dissipationless conditions), which may explain their relatively long lifetime. } {} ",
        "introduction": "Compact Groups (CG) of galaxies are local structures with high galaxy densities, equivalent to what is measured in the centers of clusters of galaxies. Contrary to clusters, however, the velocity dispersion of the galaxies in CG is low, comparable to the velocity dispersion of elliptical galaxies or the rotational velocity of spiral disks (Hickson \\cite{Hickson97} and references therein). These conditions favor interactions, and even mergers, which suggests that galaxies in CG must show plenty evidence of these phenomena, happening now, or that happened in the past. However, after almost forty years of studies, determining how CG really formed still remains an open question.  One of the most extensively studied samples of CG is the Hickson Compact Group (HCG) catalog as defined by Hickson (\\cite{Hickson82}). Subsequent spectroscopy studies have shown that the majority of the HCG have galaxies with concordant redshifts (Hickson et al. \\cite{Hickson92}), that display many signs of interactions, under the form of tidal tails, tidal bridges, distorted isophotes, faint shells, and double nuclei (Mendes de Oliveira \\& Hickson \\cite{Mendes94}; Hickson \\cite{Hickson97}; V{\\'{\\i}}lchez  \\& Iglesias-P{\\'a}ramo \\cite{Vilchez98}; Verdes-Montenegro et al. \\cite{Verdes01}, \\cite{Verdes02}, \\cite{Verdes05}; Coziol \\& Plauchu-Frayn \\cite{Coziol07}; Plauchu-Frayn \\& Coziol \\cite{Plauchu10a}). Asymmetrical rotation curves were also detected (Rubin et al. \\cite{Rubin91}; Nishiura et al. \\cite{Nishiura00}; Torres-Flores et al. \\cite{Torres10}) and, contrary to what was assumed in the past, many evidence for present and past mergers were also clearly established (Caon et al. \\cite{Caon94}; Coziol \\& Plauchu-Frayn \\cite{Coziol07}; Plauchu-Frayn \\& Coziol \\cite{Plauchu10b}). All these observations are consistent with the initial assumption which sustains that CG are physically real structures, where  galaxies during their formation and evolution have experienced frequent dynamical interactions.  Other aspects of CG, on the other hand, may still defy our comprehension. Spectroscopy study results, in particular, were found to be different than expectations based on the common assumptions that the main effects of dynamical interactions on galaxies in CG should be a direct enhancement of their star formation and activation of supermassive black holes (SMBHs) at their centers. Although emission-line galaxies are remarkably frequent in CG, involving more than 50\\% of the galaxies (Coziol et al. \\cite{Coziol98}, \\cite{Coziol00}, \\cite{Coziol04}; Mart\\'inez et al. \\cite{Martinez08}, \\cite{Martinez10}), this activity seems mostly nonthermal, under the form of Seyfert~2, LINERs, and numerous low-luminosity active galactic nuclei. Luminous AGN with typical broad line regions are relatively rare in CG (Coziol et al. \\cite{Coziol98}, \\cite{Coziol04}; Ribeiro et al. \\cite{Ribeiro96}; Mart\\'inez et al. \\cite{Martinez08}, \\cite{Martinez10}). Nuclear star formation, although mildly enhanced in some groups (e.g., Ribeiro et al. \\cite{Ribeiro96}; Iglesias-P\\'aramo \\& V\\'ilchez \\cite{Iglesias97}), is generally low (Rubin et al. \\cite{Rubin91}; Zepf et al. \\cite{Zepf91}; Moles et al. \\cite{Moles94}; Pildis et al. \\cite{Pildis95}; Menon \\cite{Menon95}; Ribeiro et al. \\cite{Ribeiro96}; Coziol et al. \\cite{Coziol98}, \\cite{Coziol00}; Verdes-Montenegro et al. \\cite{Verdes98}; Allam et al. \\cite{Allam99}; Iglesias-P{\\'a}ramo \\& V{\\'{\\i}}lchez \\cite{Iglesias99}; Mart\\'inez et al. \\cite{Martinez10}).  In part, the observation that star formation is possibly deficient and that AGN are mostly underluminous in CG could be due to a general deficiency of gas. One possibility is that this is the effect of tidal gas stripping. However, although evidence of tidal gas stripping is unquestionable in CG (e.g., Verdes-Montenegro et al. \\cite{Verdes01}; Durbala et al. \\cite{Durbala08}), this mechanism alone cannot explain all the observations. In particular, tidal gas stripping cannot explain the higher fraction of early-type galaxies found in CG (Hickson \\cite{Hickson82}; Williams \\& Rood \\cite{Williams87}; Sulentic \\cite{Sulentic87}; Hickson, Kindl \\& Huchra \\cite{Hickson88}; Palumbo et al. \\cite{Palumbo95}). Indeed, it is difficult to explain how an elliptical galaxy, or even the massive bulge of an early-type spiral galaxy, can build up after star formation is quenched in the disk of a late-type spiral galaxy by gas stripping (Dressler \\cite{Dressler80}). Unless one assumes that these galaxies can build massive bulges later by dry mergers (Plauchu-Frayn \\& Coziol \\cite{Plauchu10b}). However, neither can gas stripping explains the many AGN observed in CG (e.g., Coziol et al. \\cite{Coziol98}, \\cite{Coziol00}; Mart\\'inez et al. \\cite{Martinez10}). These AGN are related to the formation of numerous early-type galaxies through tidal triggering, a mechanism that funnels down gas right into the center of galaxies, starting a sequence of bursts of star formation that increase the mass of the bulges, and forming, or feeding, a SMBH at their centers (Merritt \\cite{Merritt83}; \\cite{Merritt84}; Byrd \\& Valtonen \\cite{Byrd90}; Henriksen \\& Byrd \\cite{Henriksen96}; Fujita \\cite{Fujita98}). However, because the level of such activities is generally low, we would have then to assume these events have taken place sometime in the recent past, or possibly during the formation of the CG (Coziol et al. \\cite{Coziol04}; Mendes de Oliveira et al. \\cite{Mendes05}). The important question is how long ago?  Recent studies using Lick indices to determine the stellar population contents of elliptical galaxies in small samples of HCG have shown that the formation process of these galaxies is similar to that of elliptical in clusters of galaxies (Proctor et al. \\cite{Proctor04}; Mendes de Oliveira et al. \\cite{Mendes05}; de la Rosa et al. \\cite{Rosa07}), and suggested their formation took place at least $2$~Gyr in the past. This is already older than predicted by the fast merger scenario for CG, which have led some authors (e.g, Proctor et al. \\cite{Proctor04}; Mendes de Oliveira et al. \\cite{Mendes05}) to propose that CG are defying the standard cosmological model of structure formation. However, this conclusion is somewhat premature, considering our current state of comprehension of the physics behind all the processes entering this model.  Many aspects of CG as described above would rather seem to be in reasonably good agreement with the basic predictions of the standard model of structure formation. In Coziol et al. (\\cite{Coziol09}) it was suggested that CG are the building blocks of larger scale structures like clusters of galaxies. Consistent with the hierarchical galaxy formation model (White \\& Frenk \\cite{White91}; Kauffmann et al. \\cite{Kauffmann93}; Cole et al. \\cite{Cole94}), galaxies form from primordial fluctuations in mass density. In regions showing the higher fluctuations, the galaxy formation process starts before than regions with lower fluctuations. Assuming the power law for the mass function of the fluctuations is relatively steep, the first structures that form would look like massive CG, where the low velocity dispersions of the protogalaxies in these systems favor interactions and mergers that produce numerous massive elliptical and early-type spiral galaxies through a sequence of starbursts. The increased potential wells of these newly assembled structures would then attract other similar systems, and a lot of primordial gas would eventually fall in, being enriched in the process by the metals expelled from galaxies by starburst winds (Torres-Papaqui et al. \\cite{TP12}). Much later, late-type spirals that formed in isolation in the field would eventually fall in, making the final structure consistent with a cluster of galaxies. According to the hierarchical galaxy formation model, in low mass density fluctuations the same process would be expected to start later, resulting, due to the small amounts of mass available, to slightly different results.   Within the above theoretical context our present knowledge of CG may suggest two different scenarios for their formation. According to the first scenario galaxies in CG have formed first in relative isolation, from very low mass density fluctuations consistent with the field, more likely forming a gas rich disk, then later on (i.e. at the present epoch), started to form a more gravitationally bound structure like the CG. Within this view, tidal gas stripping, quenching star formation in disk galaxies, and dry mergers, growing more massive bulges, would play the central role in the evolution of CG. The second scenario is more similar to what may have happened during the formation of clusters. Galaxies forming the CG never passed by a spiral form before, but merged rapidly as gas rich protogalaxy systems. Through multiple interactions, triggering starbursts and gas funneling, galaxies form massive bulges, many with an active black hole at their center. However, considering the low density of matter where these galaxies formed, and a first phase of intense star formation and AGN activity, short time-scales for gas consumption would ensue. Reaching a pure elliptical structure in these conditions would not be always possible, and the expected end products would rather be a bunch of gravitationally bounded, relatively gas deficient, early-type spiral galaxies, with an aging SMBH at their center accreting at lower rates than in the field (consistent with a LLAGN). As in the model for the formation of clusters, during the first phase of the CG formation a deeper gravitational potential well builds up, which then act as a pole of attraction for late-type galaxies that formed at the periphery, making the final structures similar to CG.  In order to distinguish between the above scenarios for the formation of CG we have used the stellar population synthesis code \\begin{scriptsize}STARLIGHT\\end{scriptsize}\\footnote{http://www.starlight.ufsc.br} (Cid Fernandes et al. \\cite{Cid05}) to determine the star formation histories (SFHs) of 210 galaxies members in 55 Hickson Compact Groups. The SFH traces the variation of star formation over the lifetime of a galaxy, and yields consequently a snapshot picture of its formation. For comparison we have also determined the SFHs of 309 galaxies taken from the Catalog of Isolated Galaxies (CIG), which represent the lowest galaxy density regime. Our study is organized in the following manner. In Sect.~\\ref{samples} we describe the selection of our two different samples. In Sect.~\\ref{method} we explain how the spectra synthesis method works. In Sect.~\\ref{results} we present the results of our comparison of the SFHs in our different samples. In Sect.~\\ref{discussion} we discuss briefly the consequence of these results and expose our main conclusions. Throughout this study a Hubble constant $H_0$ = 75 km s$^{-1}$ Mpc$^{-1}$ is assumed. ",
        "conclusions": "\\label{discussion}  \\begin{figure} \\centering \\resizebox{0.9\\hsize}{!}{\\includegraphics{ELateSpiralsSFHscaleLogByType.eps}} \\caption{Same as Fig.~\\ref{ESFHlog} for LtS galaxies. Young LtS galaxies in the samples are indicated with gray triangles.} \\label{LSSFHlog} \\end{figure}   To summarize our analysis, we show in Fig.~\\ref{AllSFHlog} the mean SFTS as measured in the three morphological classes, EtG, EtS and LtS (old subclass only), comparing the HCG  with the CIG. The most remarkable difference observed between the galaxies in the HCG (filled symbols) and CIG (open symbols) is the less prolongated star formation activity, namely SFTS, in the HCG compared to the CIG sample. The SFTS is shorter by $\\sim 2$~Gyr in the EtG and LtS, and even shorter by $\\sim 3.5$~Gyr in the EtS. These differences are independent of the morphology, mass, and luminosity of the galaxies.  For the EtG, we conclude that they have formed their stars more rapidly. In other words, our results suggest that the EtG in the HCG show higher astration rates than in the CIG (i.e., they transformed their gas more efficiently into stars). Higher astration rates for the EtG imply more frequent mergers (Coziol et al. \\cite{Coziol11}).  This is consistent with the hierarchical galaxy formation model, according to which gas rich protogalaxies in dense environments experience more frequent mergers, producing galaxies with earlier-type morphologies (i.e., galaxies with bigger bulges). However, since the CG are forming from lower mass density fluctuations than clusters of galaxies, very few of these mergers would produce massive elliptical galaxies. This picture is consistent with the morphological distribution in the HCG, as seen in Fig.~\\ref{T_DISTR}, which shows that these systems are dominated by lenticular galaxies.   For the EtS and LtS, the interpretation is not as obvious as for the EtG. In these galaxies, we cannot eliminate a priori tidal gas stripping as a possible mechanism to reduce the SFTS. These would be examples of galaxies that fell into the CG later. To test the gas stripping hypothesis, we have calculated the Specific Star Formation Rate (SSFR) of all the spiral galaxies in our samples (EtS and LtS galaxies). The SSFR, defined as the actual SFR divided by the stellar mass, reflects the mass of stars formed during the last $10^8$ yr. The SSFR is estimated using \\begin{scriptsize}STARLIGHT\\end{scriptsize} by integrating the mass-weighted population vector over the young age group $M_Y$ (see Eq. 6 in Asari et al. \\cite{Asari07}). In Fig.~\\ref{SSFRsp}a we compare the distributions of the SSFR in the HCG and  CIG samples. The distribution of SSFR is clearly bimodal (i.e. with a gap) for the HCG, while the distribution of the CIG is continuous. The bimodal distribution of the HCG is a by-product of a difference in morphology. As before, we distinguish between early and and late type spirals in both samples and find that the EtS in the HCG have a distribution of SSFR that peaks at lower values than for the CIG distribution (Figs.~\\ref{SSFRsp}b). On the other hand, the SSFR of the LtS in HCG peaks at a higher value than in the CIG sample (Fig.~\\ref{SSFRsp}c). This is consistent with the hypothesis that late-type spirals, still rich in gas, are recent acquisitions in CG. By falling into the group these galaxies increase their star formation and rapidly consume their gas (i.e., higher astration rates). Thus, by increasing their astration rates and interacting with other galaxies, these new members will transform into earlier-type spirals. Note that this interpretation do not eliminate tidal gas stripping completely, which surely may also play a role in lowering the SFTS of these galaxies. However, our test indicates that this is not the main mechanism. Recent studies, based on UV and IR data, have also found a bimodality (gap) in the distribution of SSFR in HCG galaxies and suggested that the high density environment of CGs has accelerated the evolution of galaxies transforming star forming galaxies into quiescent galaxies (Johnson et al. \\cite{Johnson07}; Tzanavaris et al. \\cite{Tzanavaris10}; Walker et al. \\cite{Walker10}; Bitsakis et al. \\cite{Bitsakis10}, \\cite{Bitsakis11}).    \\begin{figure} \\resizebox{0.97\\hsize}{!}{\\includegraphics{AllSFHscaleLogNoMassSep.eps}} \\caption{Mean SFTS for EtG, EtS, and old LtS galaxies in HCG ({\\it filled symbols}) and CIG samples ({\\it open symbols}). Error bars are $\\pm \\sigma$} of the mean values. \\label{AllSFHlog} \\end{figure}  According to our analysis, galaxies in the HCG have shorter SFTS than in the CIG because they have experienced higher astration rates. This explains why galaxies in HCG are generally deficient in gas (Williams \\& Rood \\cite{Williams87}; Verdes-Montenegro et al. \\cite{Verdes01}; Borthakur et al. \\cite{Bortha10}). The galaxies evolved more rapidly than in the field. This is consistent with the hierarchical galaxy formation model which states that CG are structures that formed recently from primordial low fluctuations in mass density. As for the time of the formation of the HCG the systematic difference in SFTS suggests this could have happened $2$ to $3.5$~Gyr in the past.  Our observations reinforce the idea that these are examples of relatively young structures as has been discussed by Hickson (\\cite{Hickson97} and references therein). However, $\\sim 3$~Gyr for the time of formation of these structures may still seem too old to explain why these systems have not collapsed into more gravitationally bound objects, like a giant blue elliptical galaxy (in fact, dynamically speaking, we should have rather expected a smaller version of a cD galaxy, i.e. an elliptical galaxy with an inflated envelope). In the introduction, we have briefly mentioned why according to the hierarchical galaxy formation model such a final state for the HCG is not plausible. This is because these structures formed from very small density fluctuation regions, where there was not enough mass to transform gas rich protogalaxies into giant elliptical or cD galaxies.    Does this mean that the HCG must be in some sort of dynamical equilibrium? This hypothesis was tested by Plauchu-Frayn \\& Coziol (\\cite{Plauchu10a}, \\cite{Plauchu10b}), who showed that galaxies in the HCG are not in an equilibrium state, but merging without gas. Now, it is worth emphasizing what is the main dynamical consequence of the absence of gas on mergers. By definition, gravity is a conservative force, which implies that to form a gravitationally bounded structure, a system must loose some of its energy. This is where the presence of free gas plays a key role. The interaction of gas-rich protogalaxies is a highly dissipative process. The gas is heated then cooled through radiation, and non isotropic gravitational interactions produce gas turbulence inducing star formation and triggering the formation of SMBH at the center of galaxies. Therefore, the first phase in the formation of galaxies in CG must be relatively fast, because of the high level of dissipation of energy produced by the gas. This is consistent with the short SFTS found in HCG and the bimodality in the SSFR distribution. However, further interactions between gas-poor galaxies (i.e., dry mergers) is a much less dissipative process, and consequently the dynamical time scales for the merger of these systems may be longer than initially assumed (a few Gyr), which would explain why we still observe CG in their present states.  To summarize, our analysis reveals that the galaxies in the HCG are older than in the CIG and that their SFTS are shorter by $\\sim 3$~Gyr. Considering that these are systematic differences, independent of the morphology, of the mass and luminosity of the galaxies, we conclude that this can only imply that the galaxies in the HCG formed their stellar population faster than the galaxies in the CIG.  We conclude that the main effect of the environment on galaxies in the HCG sample is an increase of astration rate, consistent with the fast mergers of gas-rich protogalaxies. This is consistent with the hierarchical galaxy formation model, which states that the HCG are relatively recent structures that formed from primordial low mass density fluctuations. Based on the systematic difference in SFTS we propose that the HCG most probably formed $\\sim 3$~Gyr in the past.  The galaxies in the HCG are not in equilibrium but merging without gas, which may explain why they have longer dynamical lifetime.    \\begin{figure} \\centering \\resizebox{1.0\\hsize}{!}{\\includegraphics{HistMorph.eps}} \\caption{Morphology type distribution of galaxies in the HCG and CIG samples.} \\label{T_DISTR} \\end{figure}"
    },
    "1208/1208.2496_arXiv.txt": {
        "abstract": "Although the standard $\\Lambda$CDM model describes the cosmic microwave background radiation and the large scale structure of the Universe with great success, it has some tensions with observations in the effective number of neutrino species (dark radiation) and the number of small scale structures (overabundance problem). Here we propose a scenario which can relax these tensions by producing both dark matter and dark radiation by late decays of heavy particle. Thanks to the generation mechanism, dark matters are rather warm so that the small-scale structure problem is resolved. This scenario can be naturally realized in supersymmetric axion model, in which axions produced by saxion decays provide dark radiation, while axinos from saxion decays form warm dark matter. We identify a parameter region of supersymmetric axion model satisfying all known cosmological constraints. ",
        "introduction": "The standard $\\Lambda$CDM cosmological model has been extremely successful in explaining the  observed acoustic peak in the cosmic microwave background (CMB) radiation and the formation of large scale structures (LSS). Despite its success, the $\\Lambda$CDM model seems to have some tensions with observations at small scales. The measurement of the temperature anisotropy of the CMB showed less power spectrum at small scales, suggesting that the number of effective neutrino species, $\\Neff$, has a bigger value  than the one predicted by the standard model of particle physics, so the existence of `dark radiation'. Another difficulty of $\\Lambda$CDM model is faced at small scales of the structure formation. N-body simulation with cold dark matter (CDM) has shown a tension with observation in the nonlinear regime of  structure formation, producing  more substructures in  Milky-Way galaxy size than the observed ones.  In the standard cosmological scenario, the thermal plasma after the electron-positron annihilation contains photons and neutrinos. At this epoch, total radiation energy density  can be parameterized as \\dis{ \\rho_{\\rm rad}= \\left[ 1+ \\Neff \\frac78 \\bfrac{T_\\nu}{T_\\gamma}^4 \\right]\\rho_\\gamma, \\label{Neff} } where $\\rho_\\gamma=(\\pi^2/15)T_\\gamma^4$ is the photon energy density, $T_\\nu/T_\\gamma=(4/11)^{1/3}\\simeq 1.40$ after the electron-positron annihilation, and $\\Neff$ is the effective number of neutrinos, including the contribution from dark radiation if there exists any. In the standard model with three neutrino flavors, the residual heating of the neutrino fluid due to the electron-positron annihilation slightly increases $\\Neff$, yielding   $\\Neff^{\\rm SM}= 3.046$~\\cite{Mangano:2005cc}. However, the WMAP collaboration reported  $\\Neff=4.34^{+0.86}_{-0.88} \\,(68\\% {\\rm CL})$ through  the measurements of Hubble constant and baryon acoustic oscillation~\\cite{Komatsu:2010fb}. Similarly higher values of $\\Neff$ are observed by Atacama Cosmology Telescope (ACT) and South Pole Telescope (SPT), reporting $\\Neff = 4.56\\pm 0.75$~\\cite{Dunkley:2010ge} and $\\Neff= 3.86\\pm 0.42$~\\cite{Keisler:2011aw}, respectively. It is expected that the Planck satellite will be able to measure $\\Neff$ with better precision~\\cite{Hamann:2007sb}, so make the situation more clear.   The $\\Neff$ measured in the CMB, $\\NeffCMB$, can be compared with the value $\\NeffBBN$ determined by the big bang nucleosynthesis (BBN). Observations of the primordial ${}^4{\\rm He}$ abundance provides the best constraint on $\\NeffBBN$. However there is a controversy between different groups about the relic helium abundance, e.g. \\cite{Izotov:2011pa} and \\cite{Aver:2010wq}, while another recent analysis by Mangano and Serpico~\\cite{Mangano:2011ar} gives an upper bound $\\NeffBBN \\leq 4 \\, (95\\% {\\rm CL})$. At any rate, a larger value of $\\Neff$ can be explained by extra relativistic degree of freedom existing at the epoch prior to the recombination. Many models are suggested to explain this dark radiation~\\cite{Neff models}, including the ones considering  the decays of heavy particle as the origin of dark radiation~\\cite{Neff decays}.  In the N-body simulation with CDM, the structures form hierarchically, with small structures collapsing first and merging into larger and larger bodies. CDM model describes the distribution and correlation of structures very well at large scales, however there is a large discrepancy at small scales between the observed number of satellite galaxies of the Milky Way and the expected number~\\cite{Klypin:1999uc}. This tension has brought many questions on the galaxy formation and evolution as well as the properties of dark matter. One possibility is to impose warm nature of dark matter (WDM) instead of coldness~\\cite{Bode:2000gq}. The free-streaming of WDM can reduce the power spectrum at small scales, which would  result in smaller number of galactic subhalos~\\cite{WDM}.   In fact, WDM model with $\\mw\\simeq 1 - 4\\kev$ can alleviate the CDM overabundance problems in many respects. It resolves the discrepancy in the bright satellite galaxies~\\cite{Lovell:2011rd}, solves the excess of predicted faint galaxies at low and high redshifts, as well as the excess of bright galaxies at low redshifts in the galaxy formation~\\cite{Menci:2012kk}. Also it has better agreement in the HI velocity (width) function measured in the ALFALFA survey~\\cite{Papastergis:2011xe}, and in the number of Milky Way satellites~\\cite{Polisensky:2010rw}.  In this work, we examine a scenario in which  both dark radiation and WDM find a common origin in the decays of  heavy particle. As we will see, supersymmetric axion model with relatively light saxion and axino masses provides a natural set up realizing such scenario. Decays  of massive saxion produce (nearly) massless axion pairs and massive axino pairs with different branching ratios. Axions then contribute to the dark radiation, while axinos  become warm dark matter which have large velocities to solve the small scale structure problems. ",
        "conclusions": "We have examined the possibility that late decays of massive particle after BBN can provide a common origin for dark radiation around the epoch of recombination and warm dark matter with free streaming which can solve the small scale structure problems. As a specific example, we proposed a supersymmetric axion model in which dark radiation axions and warm dark matter axinos are produced by the decays of saxion, and identified a parameter space which  can successfully realize the scenario while satisfying all the cosmological constraints."
    },
    "1208/1208.5049_arXiv.txt": {
        "abstract": "We present measurements of the auto- and cross-frequency power spectra of the cosmic infrared background (CIB) at 250, 350, and 500\\rmicron\\ (1200, 860, and 600\\,GHz) from observations totaling $\\sim 70\\, \\rm deg^2$ made with the SPIRE instrument aboard the \\emph{Herschel Space Observatory}. We measure a fractional anisotropy $\\delta I / I = 14 \\pm 4$\\%,  detecting signatures arising from the clustering of dusty star-forming galaxies in both the linear (2-halo) and non-linear (1-halo) regimes; and that the transition from the 2- to 1-halo terms, below which power originates predominantly from multiple galaxies within dark matter halos, occurs at $k_{\\theta} \\sim 0.10$--$0.12\\invarcmin$ ($\\ell \\sim 2160$--2380), from 250 to 500\\rmicron. New to this paper is clear evidence of a dependence of the Poisson and 1-halo power on the flux-cut level of masked sources --- suggesting that some fraction of the more luminous sources occupy more massive halos as satellites, or are possibly close pairs. We measure the cross-correlation power spectra between bands, finding that bands which are farthest apart are the least correlated, as well as hints of a reduction in the correlation between bands when resolved sources are more aggressively masked. In the second part of the paper we attempt to interpret the measurements in the framework of the halo model. With the aim of fitting simultaneously with one model the power spectra, number counts, and absolute CIB level in all bands, we find that this is achievable by invoking a luminosity-mass relationship, such that the luminosity-to-mass ratio peaks at a particular halo mass scale and declines towards lower and higher mass halos. Our best-fit model finds that the halo mass which is most efficient at hosting star formation in the redshift range of peak star-forming activity, $z\\sim 1-3$, is ${\\rm log}(M_{\\rm peak}/\\rm M_{\\odot}) \\sim 12.1\\pm 0.5$, and that the minimum halo mass to host infrared galaxies is ${\\rm log}(M_{\\rm min}/\\rm M_{\\odot}) \\sim 10.1\\pm 0.6$. ",
        "introduction": "\\label{sec:intro} Star formation is well traced by dust, which absorbs the UV/optical light produced by young stars in actively star-forming regions and re-emits the energy in the far-infrared/submillimeter \\citep[FIR/submm; e.g.,][]{savage1979}. Roughly half of all starlight ever produced has been reprocessed by dusty star-forming galaxies \\citep[DSFGs; e.g.,][]{hauser2001,dole2006}, and this emission is responsible for the ubiquitous cosmic infrared background \\citep[CIB;][]{puget1996,fixsen1998}. The mechanisms responsible for the presence or absence of star formation are partially dependent on the local environment (e.g., major mergers: \\citealt[][]{narayanan2010}; condensation or cold accretion: \\citealt[][]{dekel2009}, photoionization heating, supernovae, active galactic nuclei, and virial shocks: \\citealt[][]{birnboim2003, granato2004, bower2006}). Thus, the specifics of the galaxy distribution --- which can be determined statistically to high precision by measuring their clustering properties --- inform the relationship of star formation and dark matter density, and are valuable inputs for models of galaxy formation. However, measuring the clustering of DSFGs has historically proven difficult to do.  Owing to the relatively large point spread functions (PSF's) of ground-, balloon-, and space-based submillimeter observatories, coupled with very steep source counts, maps at these wavelengths are dominated by confusion noise.  For the 250\\rmicron\\  channel on \\emph{Herschel}, for example, this means that no matter how deeply you observe a field, without some sort of spatial deconvolution at best only $\\sim 15\\%$ of the flux density will be resolved into individually detected galaxies \\citep{oliver2010}. Add to that the fact that the redshift distribution of DSFGs is relatively broad \\citep[e.g.,][]{casey2012b,chapman2005,bethermin2012b}, clustering measurements of resolved sources have consequently had limited success  \\citep[e.g.,][]{blain2004, scott2006, weiss2009}, and somewhat contradictory results \\citep[e.g.,][]{cooray2010,maddox2010}.  The remaining intensities in the maps appear as fluctuations, or anisotropies, in the CIB. Contained in CIB anisotropies (or CIBA) is the clustering pattern, integrated over luminosity and redshift, of \\emph{all}  DSFGs --- including  those too faint to be resolved. And analogous to the two-point function typically used to estimate the clustering of resolved galaxies, the power spectrum of these intensity fluctuations is a probe of the clustering properties of those galaxies \\citep[e.g.,][]{bond1984,scott1999,knox2001,negrello2007}. Initial power spectrum measurements from \\emph{Spitzer} \\citep{grossan2007, lagache2007}, BLAST \\citep{viero2009, hajian2012}, ACT \\citep{dunkley2011}, and SPT \\citep{hall2010} found a signal in excess of Poisson noise originating from the clustering of DSFGs, but were limited to measuring the galaxy bias in the linear regime, rather than their distribution within dark matter halos. Subsequent measurements from \\emph{Herschel}/SPIRE \\citep{amblard2011}, and \\emph{Planck} \\citep{lagache2011} were able to isolate the linear and non-linear clustering signals, but the two groups found that their measurements agreed only after correcting for multiple systematics.  Power spectra can be interpreted with modeling frameworks in much the same way as is done for two-point function measurements of resolved sources.  Among the most commonly adopted models are so-called \\lq\\lq halo models\\rq\\rq\\ \\citep[e.g.,][]{seljak2000, cooray2002}, which use halo occupation distributions \\citep[HODs; e.g.,][]{peacock2000a,scoccimarro2001} to statistically assign galaxies to dark matter halos in order to re-create observed clustering measurements. Halo models have been adopted to interpret CIBA spectra from BLAST \\citep{viero2009}, \\emph{Herschel}/SPIRE \\citep{amblard2011, penin2012b, xia2012}, and \\emph{Planck} \\citep{lagache2011, penin2012b, shang2012, xia2012}, with varying success.  Precisely measuring the CIBA power spectra and decoding the information contained within them is a rapidly growing field, and it is also the focus of this paper. First and foremost, we aim to advance the field by providing state-of-the-art measurements of the auto- and cross-frequency power spectra of CIB anisotropies at 250, 350, and 500\\rmicron, spanning angular scales $0.01 \\le k_{\\theta} \\lsim 2\\invarcmin$ (or $350 \\lsim \\ell \\lsim 45{,}000$) (\\S~\\ref{sec:results}). With the addition of more than four times the area, we extend the efforts of \\citet{amblard2011} --- who definitively resolved a signature of non-linear clustering on small scales --- by illustrating how the strength of the non-linear clustering signal depends strongly on the flux-cut level of masked sources (\\S~\\ref{sec:clustering}). We improve on the efforts of BLAST \\citep{viero2009, hajian2012} by measuring the cross-frequency power spectra and estimate the level of correlation between bands (\\S~\\ref{sec:bandband}).  We then attempt to interpret our measurements with a series of halo models, whose common feature is to tie the luminosities of sources to their host halo masses  (\\S~\\ref{sec:formalism}), but which differ by their treatment of the spectral energy distribution (SED) of galaxy emission. Our models fit the auto and cross-frequency power spectra in each band, and measured number counts of sources, simultaneously, thereby introducing a new level of sophistication to the body of existing halo models in the literature. When required, we adopt the concordance model, a flat $\\Lambda$CDM cosmology with $\\Omega_{\\rm M} = 0.274$, $\\Omega_{\\Lambda}  = 0.726$, $H_0 = 70.5\\, \\rm km\\, s^{-1}\\, Mpc^{-1}$, and $\\sigma_8 = 0.81$ \\citep{komatsu2011}.  \\newpage ",
        "conclusions": ""
    },
    "1208/1208.2669_arXiv.txt": {
        "abstract": "The Sun is a non-equilibrium dissipative system subjected to an energy flow which originates in its core. Convective overshooting motions create temperature and velocity structures which show a temporal and spatial evolution. As a result, photospheric structures are generally considered to be the direct manifestation of convective plasma motions. The plasma flows on the photosphere govern the motion of single magnetic elements. These elements are arranged in typical patterns which are observed as a variety of multiscale magnetic patterns.\\\\ High resolution magnetograms of quiet solar surface revealed the presence of magnetic underdense regions in the solar photosphere, commonly called voids, which may be considered a signature of the underlying convective structure. The analysis of such patterns paves the way for the investigation of all turbulent convective scales from granular to global.\\\\ In order to address the question of magnetic structures driven by turbulent convection at granular and mesogranular scales we used a \u201cvoids\u201d detection method.\\\\ The computed voids distribution shows an exponential behavior at scales between 2 and 10 Mm and the absence of features at 5-10~Mm mesogranular scales. The absence of preferred scales of organization in the 2-10~Mm range supports the multiscale nature of flows on the solar surface and the absence of a mesogranular convective scale. ",
        "introduction": "\\label{S-Introduction}  The Sun is a non-equilibrium dissipative system subjected to an energy flow which originates in its core. A turbulent convective envelope is generated by this energy flow in response to the significant opacity of the outer solar region. The photosphere shows a state of reduced spatial symmetry due to the convective overshooting motions on the surface. The result of this reduced symmetry is the formation of temperature and velocity structures which show a temporal and spatial multiscale evolution. As a result, photospheric structures are generally considered to be the direct manifestation of convective motions.\\\\ Global plasma flows (i.e., differential rotation, meridional circulation and torsional oscillation) and the solar magnetic field, interacting with these turbulent convective motions, increase the complexity of emerging structures and of the star as a whole. The dynamics of the solar surface and its interaction with the magnetic field ultimately control the  structure of the outer solar atmosphere and the heliosphere beyond.\\\\ The plasma flows on the photosphere govern the motion of single magnetic elements. These elements are arranged in typical patterns which are observed as a variety of multiscale magnetic features, e.g., filamentary clusters or clumps of magnetic elements, magnetic network, clusters of facular points, active regions. The analysis of these multiscale magnetic patterns on the solar surface paves the way for the investigation of all turbulent convective scales from granular to global. \\opencite{Chaouche} report, indeed, that {\\it the magnetic field can provide a direct avenue to explore convective patterns since the magnetic flux can be measured directly using Stokes polarimetry techniques}. \\\\ Commonly, four different convective scales are identified on the solar surface: the granulation, which shows a typical horizontal scale of about 1 Mm and a short (a few minutes) lifetime, the mesogranulation, which shows horizontal length scales ranging from 5 Mm to 10 Mm and a lifetime of some hours, the supergranulation, which shows typical horizontal length scales from 20 Mm to 50 Mm and a lifetime of about one day, and giant cells which show horizontal length scales of the order of 100Mm or larger.\\\\ The wide range of values which are reported in the observed horizontal length scales of mesogranulation and supergranulation are most likely due to different methods of data analysis and choices of spatial and time filters. Even more important, however, is that this division, as suggested in \\opencite{Nordlund2009}, {\\it is largely of historical origin, and current evidence indicates that there is a continuous spectrum of motions, on all scales from global to sub-granular}.\\\\  In the last years, high resolution magnetograms of quiet solar surface (e.g., \\opencite{berger}, \\opencite{almeida}, \\opencite{lites08}) revealed the presence of multiscale magnetic underdense regions in the solar photosphere, commonly called voids, which may be considered a signature of the underlying convective structure.  From magneto-convection simulations it results that at the granular scales the magnetic field elements are concentrated at the boundaries of granular cells (i.e., intergranular lanes). Several observations, i.e., \\opencite{cerdena03a}, \\opencite{Khomenko03}, \\opencite{manso}, confirmed such numerical results. The presence of magneto-convective concentration at mesogranular scales has been predicted by numerical simulations (\\opencite{Cattaneo01}) and observed with 2D spectro-polarimetry (\\opencite{cerdena03b}, \\opencite{Chaouche}), IR spectropolarimetry (\\opencite{truillo03}) and from magnetograms (\\opencite{almeida}).\\\\ In more detail, \\opencite{cerdena03b} assessed the existence of a web-like pattern with the spatial scale of mesogranulation by using a time sequence of Inter-Network magnetograms observed at the solar disk center. The authors reported that the observed persistent pattern resembled a network with a spatial scale between 5 and 10 arcsec, which they identified as mesogranulation.  More recently, \\opencite{Chaouche} confirmed the preferential location of magnetic elements in mesogranular cells using high spatial and temporal resolution IMaX-SUNRISE dataset. The mesogranular cells were identified using Lagrange tracers driven by horizontal velocity fields which were computed via Local Correlation Tracking technique (\\opencite{November88}). In addition, the statistical multiscale analysis of the photospheric field geometry outlined the role of solar surface convective motions in the clustering and intermittent organization of magnetic flux elements in the quiet Sun (\\opencite{lawrence93},\\opencite{Uritsky}). Moreover, \\opencite{deWijnM} reported that it seems highly likely that the cells described by \\opencite{deWijn05} and the \u201cvoids\u201d found by \\opencite{lites08} correspond to mesogranules.\\\\  In order to address the question of which organization scales appear in magnetic patterns formed by turbulent convective motions and how these scales could be revealed in high resolution magnetograms, we used a \u201cvoid\u201d detection method. Void detection methods are largely used in cosmology to study the voids from galaxy and clusters (\\opencite{aikio98}). These numerical methods define void structures of different particle distributions. In these methods galaxies are considered single {\\it particles}. High-density particle regions surrounding low-density regions actually form the edges of these empty volumes, called voids. In the case of magnetic patterns on the solar surface we observe areas of low or absent magnetic activity surrounded by intense magnetic structures. Like the voids in the distribution of galaxies we define this underdense magnetic field regions in solar magnetograms as voids.\\\\  In this work we employ a suitable automated void detection procedure to single out these underdense regions and to investigate their geometrical properties. We concentrate our analysis on the remarkable scales of granulation, the only length scale for which the convective signature is directly visible on the solar surface, and mesogranulation, whose physical origin is still under debate. See \\opencite{Nordlund2009} for a recent review on solar surface convection. In more detail, we are concerned with the study of distribution of voids in a quiet Sun high resolution magnetogram observed at disk center.\\\\  \\begin{figure}[h!] \\includegraphics[width=1\\textwidth,clip=]{magnetogram.eps} \\caption{Hinode SOT/SP line-of-sight magnetogram of a $302 \\times 162$ $arcsec^{2}$ portion of the solar photosphere observed at disk center. } \\label{magnetogram} \\end{figure} ",
        "conclusions": "\\label{S-conclusions}  As reported in paragraph \\ref{S-algorithm}, we applied the algorithm to a large SOT-HINODE magnetogram and the result is that 1951 magnetic underdense regions, i.e., voids, were identified.\\\\ The calculated  distribution of these 1951 voids (Fig.\\ref{distributionHINODE}) shows an exponential decrease between 2 and 10 Mm with a decay constant equal to $2.2\\pm0.2$ Mm and a coefficient of determination $R^2=0.98$. The magnitude of $R^2$ indicates a very high degree of correlation. The absence of features at scales of 5-10~Mm indicates the lack of an intrinsic mesogranular scale.\\\\ The effect of different threshold values is investigated from a comparison of corresponding void distributions. Our findings show that the decay constant decreases when the threshold increases from 120~Gauss to 260~Gauss. The distributions remain exponential with the same high degree of correlation, even though the mesogranular signature is not observed.\\\\ Thus, we support the results reported by \\opencite{Chaouche} that mesogranulation is not among the primary energy-injection scales of solar convection. In addition, these authors report that the Probability Density Function of the distance between magnetic footpoints shows a constancy of the slope at scales between 1 and 10 Mm and that the characteristic decay distance is approximately 1.7 Mm. Moreover, they reported that only after about 20 minutes of integration a sharp mesogranular network appeared.\\\\ Similarly, \\opencite{berrilli} found that mesogranular photospheric features appear on the solar surface with a characteristic growth time of about 10~minutes. They used the hexagonal normalized information entropy measure, $H'_{hex}$, to observe the mesogranular scale emergence and to compute the associated time of formation. They found that the $H'_{hex}$ signal shows a maximum on the mesogranular scale of 8~Mm after 15-20 minutes of integration.\\\\ \\begin{figure}[h!] \\centering\\includegraphics[width=0.65\\textwidth,clip=]{hinode_histo.eps} \\caption{Void frequency distribution of 1951 identified voids.} \\label{distributionHINODE} \\end{figure} \\begin{figure}[h!] \\centering\\includegraphics[width=0.65\\textwidth,clip=]{simul_MG_2012_100.eps} \\caption{Average void distributions for pseudorandom patterns of structures (dot) and pseudorandom patterns of structures with 100 round empty regions which generate voids of mesogranular scale (heavy continuum). In the second case, the final distribution shows two populations. The first population is identified by the 3~Mm maximum related to the correlation distance among structures. The second population is identified by the second maximum related to the mesogranular scale voids.} \\label{distributionS} \\end{figure}  The absence of specific scales of organization in the 2-10~Mm range supports the multiscale nature of flows on the solar surface, but does not rule out the possibility of convection structures with these dimensions. For example, large scale flows can be obtained as the result of the local merge of downward plumes in advective-interaction models (\\opencite{rast}).\\\\ Exponential functions describe other physical quantities in turbulent convective systems, e.g., the probability density of normalized temperature fluctuations shows an interesting feature at the transition from soft to hard turbulent convection (\\opencite{yakhot}). As stated by authors, in the soft turbulent convective regime, when the Rayleigh number $Ra < 10^{7}$, the probability distribution measured at the center of the cell is Gaussian, while in the hard-turbulence regime ($Ra>10^{8}$) the probability distribution is very close to exponential. The value of the Rayleigh number characterizing intensity of the turbulent convection in photosphere is $Ra\\simeq10^{11}$ (\\opencite{Bray}) which is consistent with the hard-turbulence regime.\\\\  The hypothesis that there is a continuous spectrum of flows on the solar surface, from global to sub-granular (\\opencite{Nordlund2009}), still needs to be thoroughly \\begin{tiny} \u2022 \\end{tiny}proven. With this in mind, we are developing a novel fast void detection algorithm to analyze a large dataset of MDI magnetograms acquired during the \u201cexceptional\u201d solar minimum at the end of solar cycle 23. The analysis of this dataset will allow us to investigate voids in the range of 5-100~Mm in order to validate the above hypothesis.   \\begin{acks} We wish to thank Bartolomeo Viticchi\\`e for his contribution in the analysis of the HINODE data and Michael Senno and Roberto Piazzesi for their contributions in reviewing the manuscript. This project is supported by the University of Rome Tor Vergata Astronomy Ph.D. Program. Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in cooperation with ESA and NSC (Norway). \\end{acks}   \\mbox{}~\\\\"
    },
    "1208/1208.2033_arXiv.txt": {
        "abstract": "{Gamma-ray observations provide sensitive tests of Lorentz invariance violation (LIV). At present the most sensitive tests come from observations of transient events, gamma-ray bursts and flaring AGN. Disadvantages of transients are that an independent confirmation by a different experiment is often not possible and that limits cannot be improved with a longer exposure. Pulsars do not have these disadvantages. Testing Lorentz invariance with pulsars was not considered seriously so far because limits were not competitive. The VERITAS collaboration has recently reported the detection of pulsed emission from the Crab pulsar above 100\\,GeV. This measurement can be used to constrain LIV effects with a sensitivity that is competitive with some of the best available limits. In view of this unexpected result we discuss what the prospects are of doing LIV tests with very-high energy gamma-ray emission from pulsars. } ",
        "introduction": "Lorentz invariance (LI) is a fundamental concept of modern physics. One of the consequences of LI is that the speed of light is constant and, in particular, that it does not depend on the photon energy. If, however, spacetime has structure, like it is postulated in some approaches to combine quantum mechanics and general relativity \\cite{Smolin, Amelino:98}, it could be that the speed of light depends on the energy of the photon. From our everyday experience we know that the effect, if it exists, is very small and, therefore, only evident if the energy of the photons is in the gamma-ray regime and the photons have traveled large distances. In conclusion, Lorentz invariance violation (LIV) can be tested effectively by searching for a delay in the arrival time of gamma rays with different energies that have been emitted simultaneously from an astrophysical object.  A quantitative prediction of the energy dependency of the speed of light does not exist. A pragmatic ansatz that is often made is to modify the constant speed of light $c$ by adding energy dependent terms that are proportional to $(E/E_{\\rm LIV})^n$, where $E$ is the energy of the photon, and $E_{\\rm LIV}$ can be considered as the energy scale at which Lorentz invariance violating effects become evident. $E_{\\rm LIV}$ is the quantity that is normally tested in LIV tests. With this ansatz, two photons with different energies $E_h$ and $E_l$ that are emitted simultaneously arrive at an observer at a distance $d$ at slightly different times. Depending on the order $n$ of the energy dependence, the time difference $\\Delta t$ for the linear and the quadratic term in $E/E_{\\rm LIV}$ are: \\begin{equation}\\label{lin} \\Delta t_1 = \\frac{d}{c} \\cdot\\frac{E_{\\rm h} - E_{\\rm l}} {E_{\\rm LIV}}\\rightarrow E_{\\rm LIV} =  \\frac{d}{c} \\cdot\\frac{E_{\\rm h} - E_{\\rm l}} {\\Delta t_1} \\end{equation} and \\begin{equation}\\label{quad} \\Delta t_2 = \\frac{d}{c}\\cdot \\frac{3}{2} \\cdot \\frac{E^2_{\\rm h} - E^2_{\\rm l}}  {E^2_{\\rm LIV}}\\rightarrow E_{\\rm LIV} =\\sqrt{\\frac{d}{c}\\cdot \\frac{3}{2} \\cdot \\frac{E^2_{\\rm h} - E^2_{\\rm l}}  {\\Delta t_2}}\\,, \\end{equation} respectively.  In Section 2 I discuss the status of LIV tests with flaring AGN and GRB, which provide the most stringent tests of $E_{\\rm LIV}$ to date. In Section 3 the advantages of doing LIV tests with pulsars are explained. In Section 4 I give a quantitative estimate of how much LIV can be constrained with the recent detection of the Crab pulsar at 120\\,GeV by the VERITAS collaboration \\cite{andrew}. In Section 5 I discuss the prospects of doing LIV tests with pulsars. The paper closes with summarizing remarks in Section 6. ",
        "conclusions": "The detection of the Crab pulsar above 100\\,GeV shows that LIV tests with pulsars are competitive with other LIV tests. A robust estimate shows that limits on LIV can be derived with the VERITAS Crab pulsar detection that are one order of magnitude below the best available limit from an AGN observation. Pulsars also provide a unique way of separating source intrinsic effects from propagation effects, e.g.\\ LIV. The limits can be improved in the future with more refined analysis methods and deeper observations. With CTA it should be possible to probe LIV with pulsar observations at the Planck mass."
    },
    "1208/1208.6574_arXiv.txt": {
        "abstract": "{ The energy spectrum of ultra-high energy cosmic rays above 10$^{18}$~eV is measured using the hybrid events collected by the Pierre Auger Observatory between November 2005 and September 2010. The large exposure of the Observatory allows the measurement of the main features of the energy spectrum with high statistics. Full Monte Carlo simulations of the extensive air showers (based on the CORSIKA code) and of the hybrid detector response are adopted here as an independent cross check of the standard analysis~\\cite{spectrum2010}. The dependence on mass composition and other systematic uncertainties are discussed in detail and, in the full Monte Carlo approach, a region of confidence for flux measurements is defined when all the uncertainties are taken into account. An update is also reported of the energy spectrum obtained by combining the hybrid spectrum and that measured using the surface detector array. \\PACS{ {96.50.S-}{Cosmic rays}   \\and {96.50.sb}{energy spectra} \\and {96.50.sd}{extensive air shower} \\and {98.70.Sa}{galactic and extragalactic} } % } % \\authorrunning{M. Settimo for the Pierre Auger Collaboration} \\titlerunning{Cosmic ray energy spectrum with the Pierre Auger Observatory} ",
        "introduction": "The features of the energy spectrum of ultra-high energy cosmic rays are intrinsically connected to the origin, nature and propagation of cosmic rays. At the highest energies, above 4$\\times$10$^{19}$~eV, a suppression of the flux has been observed by the HiRes experiment~\\cite{hires_spectrum}, the Pierre Auger Observatory~\\cite{spectrum2010,SDspectrum,spectrum2011} and the Telescope Array~\\cite{TA_spectrum}. This suppression is compatible with the predicted Greisen-Zatsepin-Kuz'min (GZK) effect~\\cite{gzk1,gzk2}, even if other possibilities (\\emph{e.g.} limits in the maximum energy at the source) cannot be excluded. A break in the power law spectrum, named  the ``ankle'', has also been reported around 10$^{18.6}$~eV~\\cite{spectrum2010,hires_spectrum,spectrum2011,TA_spectrum,linsley,ankleexp,akeno,flyseye}. This feature is traditionally explained as the intersection of a steep Galactic component~\\cite{linsley,hillas,ankle1,ankle2,ankle3} with a flatter extragalactic one though in this case the galactic component must extend up to energies above 10$^{18}$~eV, requiring a modification of the simple rigidity model of the cosmic ray confinement in the galaxy~\\cite{knee}. Other models explain the ankle structure as the distortion of a proton dominated spectrum through e$^{+}$/e$^{-}$ pair production of protons with the photons of the cosmic microwave background~\\cite{hillas2,blumenthal,dip1,dip2}. A measurement of the cosmic ray flux in this energy range together with the knowledge of the mass composition over a wide energy range~\\cite{massAuger1,massAuger2,massHiRes,massHiResMIA,massYakutsk1,massYakutsk2,massTA} may help to confirm these results and constrain different model scenarios.  The energy spectrum above 2.5~$\\times$10$^{18}$~eV has been derived using data from the surface detector array of the Pierre Auger Observatory~\\cite{SDspectrum}. This measurement has been extended to 10$^{18}$~eV~\\cite{spectrum2010,spectrum2011} using the hybrid events, that are simultaneously observed by the fluorescence telescopes and by the surface detector.  In this paper the measurement of the energy spectrum is updated to September 2010 and a method using detailed simulations of the extensive air showers (EAS) and of the hybrid detector response has been also developed. Hereafter we refer to this approach as ``full Monte Carlo''. It provides a complete treatment of the shower-to-shower fluctuations and an independent validation of the standard method (``fast simulations'') used in~\\cite{spectrum2010} and described in detail in~\\cite{exposure}. The standard method allows one to simulate a huge amount of events and to apply stricter analysis cuts which reduce the systematic uncertainties on the spectrum measurement. The two approaches adopted in this paper differ in the EAS and detector simulations and in the selection of events. Their advantages and drawbacks are discussed and the systematic uncertainties related to mass composition, hadronic interaction models and efficiency of the detector are studied.  The paper is organized as follows. In section~\\ref{sect:auger} we introduce the Pierre Auger Observatory~\\cite{augerNIM} and the hybrid detection mode, discussing its performance. The successive sections describe the steps to derive the energy spectrum, including the on-time and the hybrid exposure calculation using both the full Monte Carlo and the fast simulation approaches. The latter is introduced to provide larger simulation statistics in combination with a more strict analysis aiming to reduce the systematic uncertainties on the energy spectrum derived in section~\\ref{sect:spectrum}. ",
        "conclusions": "The measurement of the cosmic ray flux above 10$^{18}$~eV has been updated to September 2010 using hybrid events of the Pierre Auger Observatory. The standard approach used here, and already adopted in a previous publication~\\cite{spectrum2010}, is based on fast CONEX and detector simulations. In this paper the energy spectrum has additionally been derived using a full Monte Carlo method, based on CORSIKA air showers and detailed simulations of the hybrid detector. The full Monte Carlo approach provides a complete treatment of the shower-to-shower fluctuations, even in a region where the detector is not fully efficient and is an independent validation of the standard method. Producing a huge number of events is however computationally demanding. The lack of accurate knowledge of the mass composition propagates to the spectrum giving a confidence region for the expected flux. This is defined by the two extreme assumptions of pure proton and pure iron composition. Tighter cuts, designed to reduce this systematic uncertainty, are used in the standard method, profiting from the enormous statistics provided by the fast simulations. The average values of the spectra derived with the two approaches agree to within a few percent. In both cases, the dominant contribution to the systematic uncertainties in the flux measurement comes from the overall uncertainty on the energy scale, which is estimated to be 22\\%. The energy spectrum from the standard approach has been combined with the one derived above 10$^{18.5}$~eV by the surface array between January 2004 and December 2010. This updated combination provides an accurate determination of the spectral features in the energy range between 10$^{18}$~eV and 10$^{20}$~eV. The position of the ankle has been found to be at log$_{10}(\\rm{E/eV})\\,=\\,18.61\\,\\pm\\,0.01$ and a flux suppression has been observed at log$_{10}(\\rm{E/eV})\\,=\\,19.63\\,\\pm\\,0.02$, with a significance larger than 20~$\\sigma$."
    },
    "1208/1208.3351_arXiv.txt": {
        "abstract": "We describe the motivation, design and implementation of the CORNISH survey, an arcsecond resolution radio continuum survey of the inner Galactic plane at 5\\,GHz using the Karl G. Jansky Very Large Array (VLA). It is a blind survey co-ordinated with the northern {\\it Spitzer} GLIMPSE I region covering 10\\degs$<l<$65\\degs\\ and $|b|<$1\\degs\\ at similar resolution. We discuss in detail the strategy that we employed to control the shape of the synthesised beam across this survey that covers a wide range of fairly low declinations. Two snapshots separated by 4 hours in hour angle kept the beam elongation to less that 1.5 over 75\\,\\% of the survey area and less than 2 over 98\\,\\% of the survey. The prime scientific motivation is to provide an unbiased survey for ultra-compact $\\hii$~regions to study this key phase in massive star formation.  A sensitivity around 2\\,mJy will allow the automatic distinction between radio loud and quiet mid-IR sources found in the {\\it Spitzer} surveys. This survey has many legacy applications beyond star formation including evolved stars, active stars and binaries, and extragalactic sources. The CORNISH survey for compact ionized sources complements other Galactic plane surveys that target diffuse and non-thermal sources as well as atomic and molecular phases to build up a complete picture of the ISM in the Galaxy. ",
        "introduction": "The study of star formation, stellar populations and Galactic structure is currently being transformed by a new generation of Galactic plane surveys. These surveys have much higher resolution and sensitivity than previous surveys and are at infrared and longer wavelengths that can penetrate the high extinction of the Galactic mid-plane. They also cover sufficiently wide areas of the plane to provide samples that are large enough to be representative and statistically robust for rare and short-lived phases of evolution.  Blind surveys that cover contiguous, wide areas also allow the study of phenomena and trends across a large range of scales.  A wide wavelength coverage is required to distinguish and characterise the different populations found at both the early and later phases of stellar evolution and so it is also important to ensure that the same wide areas are covered by all the relevant wavelengths. Here we describe the design and implementation of a new radio continuum survey of the Galactic plane and its role in the context of the other multi-wavelength surveys.  \\subsection{A New Generation of Galactic Plane Surveys}  Advances in various wide-field technologies have led to a new generation of Galactic plane surveys.  Those with relevance here, and to star formation in particular, are listed in Table~\\ref{surveys}. Leading the way is the GLIMPSE I (Galactic Legacy Infrared Mid-Plane Survey Extraordinaire) survey (Benjamin et al. 2003; Churchwell et al. 2009). This {\\it Spitzer} IRAC Legacy Programme covered the inner galaxy ($10\\degs<l<65\\degs$ and $-65\\degs<l<-10\\degs$, $|b|<1\\degs$) at 3.6\\,\\mic, 4.5\\,\\mic, 5.8\\,\\mic\\ and 8.0\\,\\mic. GLIMPSE is two orders of magnitude more sensitive (limit at L-band (3.6\\,\\mic)\\,$\\approx14-15$ magnitude although usually confusion limited) and has ten times higher spatial resolution ($1.4-1.9\\arcsec$) than any previous mid-IR survey. It has catalogued over 49 million sources in a wavelength regime which preferentially selects sources with hot circumstellar dust emission such as young and evolved stars (e.g., Robitaille et al. 2008). More embedded and cooler sources are the subject of the MIPSGAL survey with {\\it Spitzer} (Carey et al. 2009) and the PACS element of the Hi-GAL survey with Herschel (Molinari et al. 2010).  \\begin{deluxetable}{@{}l@{}c@{}c@{}c@{~}c@{~}c@{~}l@{}} \\tablecaption{Galactic plane surveys showing the context of the CORNISH survey. \\label{surveys}} \\tablewidth{0pt} \\tablehead{ \\colhead{Survey} & \\colhead{Wavelength} & \\colhead{Beam (\\arcsec)} & \\colhead{$l$ Coverage}& \\colhead{$b$ Coverage} & \\colhead{Probe} & \\colhead{Reference}} \\startdata IPHAS & H$\\alpha$  & 1.7 & $30$\\degs $<l<210$\\degs & $|b|<5$\\degs & Nebulae \\& stars  & Drew et al. (2005) \\\\ UKIDSS & JHK & 0.8 &  $-2$\\degs $<l<230$\\degs & $|b|<1$\\degs &  Stars, Nebulae & Lucas et al. (2008) \\\\ VVV & ZYJHK & 0.8 & $-65$\\degs $<l<10$\\degs & $|b|<2$\\degs & `` & Minniti et al. (2010)  \\\\ GLIMPSE & 4-8\\,\\mic & 2 &  $-65$\\degs $<l<65$\\degs  &$|b|<1$\\degs & Stars, Hot Dust & Churchwell et al. (2009) \\\\ MSX & 8-21\\,\\mic & 18 & All &$|b|<5$\\degs & Warm Dust & Price et al. (2001) \\\\ MIPSGAL & 24,70\\,\\mic & 6, 20 & $-65$\\degs $<l<65$\\degs &$|b|<1$\\degs & `` & Carey et al. (2009) \\\\ AKARI & 50-200\\,\\mic & 30-50 & All sky & & Cool Dust & White et al. (2009) \\\\ Hi-GAL & 70-500\\,\\mic & 10-34 &  All  & $|b|<1$\\degs\\tablenotemark{a} & `` & Molinari et al. (2010) \\\\ JPS   & 450,850\\,\\mic & 8-14 & $10$\\degs $<l<60$\\degs &$|b|<1$\\degs & `` & Moore et al. (2005) \\\\ ATLASGAL & 850\\,\\mic & 19 &$-60$\\degs $<l<60$\\degs &$|b|<1.5$\\degs & `` & Schuller et al. (2009) \\\\ BOLOCAM & 1100\\,\\mic & 33 & $-10$\\degs $<l<90$\\degs & $|b|<0.5$\\degs& `` & Aguirre et al. (2010) \\\\ GRS  & $^{13}$CO 1-0 & 46 & $18$\\degs $<l<56$\\degs &$|b|<1$\\degs& Molecular Gas & Jackson et al. (2006) \\\\ MMB & 6.7\\,GHz & 192\\tablenotemark{b} & $-180$\\degs $<l<60$\\degs &$|b|<2$\\degs & Methanol Masers & Green et al. (2009) \\\\ HOPS & 22\\,GHz & 132\\tablenotemark{b} & $-180$\\degs $<l<60$\\degs &$|b|<2$\\degs & Water Masers & Walsh et al. (2011) \\\\ OH & 1.6\\,GHz & $\\sim$10 &  $-45$\\degs $<l<45$\\degs & $|b|<3$\\degs & Hydroxyl Masers & Sevenster et al. (2001) \\\\ CORNISH & 6 cm & 1.5 &  $10$\\degs $<l<65$\\degs &$|b|<1$\\degs &Compact Ionized Gas & This work \\\\ S/V/CGPS & 21 cm & 60 &  $-107$\\degs $<l<147$\\degs & $|b|<1.3$\\degs &Atomic Gas & Stil et al. (2006) \\\\ MAGPIS & 20 cm & 5 &  $5$\\degs $<l<48$\\degs & $|b|<0.8$\\degs &Diffuse Ionized Gas & Helfand et al. (2006) \\\\ MGPS-2  & 35 cm & 45 &  $-115$\\degs $<l<0$\\degs &$|b|<10$\\degs & `` & Murphy et al. (2007) \\\\ \\enddata \\tablenotetext{a}{Follows warp} \\tablenotetext{b}{Maser positions are accurate to about 0.1\\arcsec\\ after interferometric follow-up.} \\end{deluxetable}  An even deeper probe of the general stellar population is becoming available via the deep near-IR Galactic Plane Survey (GPS) that is part of the UK IR Deep Sky Surveys (UKIDSS) programme (Lucas et al. 2008).  This survey is about 3 magnitudes more sensitive and two to three times better resolution than the 2MASS all sky near-IR survey (Skrutskie et al. 2006) and will detect upwards of a billion sources. The addition of near-IR data, where photospheric and scattered light contributions dominate, allows a much better separation and characterisation of the general stellar population, than from mid-IR colours alone. Near-IR surveys are also being used to map the extinction due to molecular clouds (e.g., Rowles \\& Froebrich 2009). The optical H$\\alpha$ survey IPHAS (Drew et al. 2005) is providing both stellar characterisation, 3D extinction maps and tracing nebular emission in less obscured regions of the northern plane with the southern VPHAS+ survey (see www.vphasplus.org) scheduled to do the same in the south.  To place the stellar populations in a Galactic context, commensurate studies of the various components of the ISM are required. For the study of star formation in particular, the molecular component is provided in part by the BU-FCRAO $^{13}$CO\\,1-0 Galactic Ring Survey (Jackson et al. 2006), which covers over half of the northern GLIMPSE area with extensions over much of the northern mid-plane (Mottram \\& Brunt 2010). The mostly optically thin $^{13}$CO\\,1-0 data traces the distribution and dynamics of the cold molecular gas. The cool dust in these molecular clouds can also be mapped via its sub-millimetre continuum emission. This is currently being achieved with the ATLASGAL survey (Schuller et al. 2009) at 870\\,\\mic\\ and BOLOCAM survey at 1100\\,\\mic\\ (Aguirre et al. 2010), whilst the forthcoming JCMT Galactic Plane survey (JPS, www.jach.hawaii.edu/JCMT/surveys/) will map the continuum at 450 and 850\\,\\mic\\ with higher sensitivity and resolution. Together with the SPIRE element of the Hi-GAL survey with the Herschel satellite (Molinari et al. 2010) this will provide temperature and dust emissivity information across the GLIMPSE region.  The more widely distributed atomic hydrogen component is covered by the International Galactic Plane Survey at about 1\\arcmin\\ resolution (www.ras.ucalgary.ca/IGPS/). This is made up of the VLA Galactic Plane Survey (VGPS, Stil et al. 2006), Southern Galactic Plane Survey (SGPS; McClure-Griffiths et al. 2005) and Canadian Galactic Plane Survey (Taylor et al. 2003). A sub-section of the plane is covered by the much deeper GALFA survey of H\\,{\\scriptsize I} at 4\\arcmin\\ resolution using the Arecibo telescope (Peek et al. 2011). The H\\,{\\scriptsize I} component traces where most of the mass of interstellar gas is located and its velocity structure.  21\\,cm absorption lines can also help to solve the near/far distance ambiguity in relation to the kinematic distances derived towards the embedded $\\hii$~regions (Sewilo et al. 2004; Fish et al. 2003)  and dense, molecular clouds (Jackson et al. 2002; Gibson et al. 2005; Busfield et al. 2006; Roman-Duval et al. 2009). This combination of atomic and molecular kinematic data can therefore provide the 3D Galactic setting of the neutral gas.  The one major component missing from the multi-wavelength surveys of the Galaxy is that of the ionized gas. No existing radio continuum survey has sufficient resolution, depth or coverage to provide the data that would complete the picture of the Galaxy. Ionized gas arises in nebulae and stellar winds around hot stars and is crucial to understanding key phases of the early and late evolution of high and intermediate mass stars.  Dense, photo-ionized regions around young and evolved stars are often optically thick with a turnover frequency of around a few GHz, where the radio continuum spectral index $\\alpha$ ($S_{\\nu}\\propto\\nu^{\\alpha}$) changes from the optically thick value of +2.0 to the optically thin slope of $-0.1$. Ionized stellar winds have a positive spectral index $\\alpha\\sim+0.6$, and so it is also much more efficient to carry out systematic searches for these sources at high frequencies. Furthermore, the resolution should be comparable to the arcsecond resolution of the IR surveys to allow IR counterparts to be uniquely identified. It was to fill this need for a high frequency, high resolution radio continuum survey covering the area of the {\\it Spitzer} survey that the concept of the Co-Ordinated Radio `N' Infrared Survey for High-mass star formation (CORNISH\\footnote{This acronym derives from the Celtic origins of Hoare and Diamond.}) survey was conceived.  \\subsection{Previous Radio Surveys}  Many single-dish surveys of the Galactic plane have been carried out at 5\\,GHz with a resolution of a few arcminutes that is totally inadequate to complement the new generation of IR surveys (e.g., Altenhoff et al. 1978; Haynes et al. 1978; Gregory \\& Condon 1991; Griffith et al. 1994). Most higher resolution interferometric surveys have been carried out at the relatively low frequency of 1.4 GHz (Zoonematkermani et al. 1990; Becker et al. 1990; Condon et al. 1998; Giveon et al. 2005).  Although the multi-configuration MAGPIS survey (Helfand et al. 2006) and MGPS-2 survey (Murphy et al. 2007) are good for tracing the optically thin thermal emission from extended, evolved $\\hii$~regions and non-thermal supernova remnants, they are not appropriate for the study of dense, thermal sources due to the unfavourable spectral index and compactness.  Becker et al. (1994), with infills by Giveon et al. (2005) and White et al. (2005) have surveyed part of the inner galaxy ($-10\\degs<l<42\\degs$, $|b|<0.4\\degs$) at 5\\,GHz with the VLA in C configuration giving a resolution of 4\\arcsec$\\times$9\\arcsec. However, these surveys cover only one fifth of the northern GLIMPSE region and the spatial resolution is poorer than that of {\\it Spitzer} IRAC, which causes problems with source identification. At this resolution the vast majority of their sources are unresolved and thus we have no morphological information. In the southern hemisphere the AT20G survey at 20\\,GHz and $\\sim$30\\arcsec\\ resolution (Murphy et al. 2010) that was primarily designed to study the extragalactic radio source population is detecting many $\\hii$~regions, but again not at sufficient depth or resolution to address the science opened up by current surveys of the Galactic plane at infrared wavelengths. ",
        "conclusions": "The CORNISH survey is a new sensitive, high spatial resolution Galactic plane survey at 5\\,GHz. It is targeted at compact thermal sources and $\\uchii$ regions in particular.  Conceived to complement the {\\it Spitzer} GLIMPSE survey, it will systematically address key questions in massive star formation such as the lifetime and evolution of the $\\uchii$ phase as a function of exciting star parameters and environment. Uniform coverage of the northern GLIMPSE region enables the distinction between radio-loud and radio-quiet objects such as MYSOs/$\\uchii$s and PPN/PN that otherwise have very similar mid-IR colours.  Legacy survey science in combination with other recent and upcoming optical, near-IR, far-IR, sub-mm and longer wavelength radio surveys of the plane will be possible with applications in evolved stars, active stars and active binaries.  Data taking for the CORNISH survey has been completed and the basic goal of uniformity of coverage achieved. The data reduction, source extraction and statistical properties of the survey and sources will be discussed in the forthcoming paper by Purcell et al. (2012). At that time the catalogue and image data will be publically available at http://www.ast.leeds.ac.uk/cornish.  We have also achieved the aim of controlling the beam shape across the survey such that the elongation of the synthesised beam is mostly less than 1.5 and nearly always less than 2. This was achieved by using a strategy of two snapshots separated by 4 hours in hour angle. A more uniform and rounder beam could be achieved by more snapshots but that would increase the overheads and has diminishing returns with the VLA. Our simulations show that snapshots separated by 6~hours in hour angle would have produced a better beam. However, this would have complicated the scheduling of the survey into 8~hour runs and more than likely compromised the very uniform coverage that we ultimately achieved. Such tradeoffs will be the case for any survey with the EVLA at these declinations and for future Galactic plane surveys in particular."
    },
    "1208/1208.3217_arXiv.txt": {
        "abstract": "Supernovae (SNe), the luminous explosions of stars, were observed since antiquity, with typical peak luminosity not exceeding $1.2 \\times 10^{43}$\\,erg\\,s$^{-1}$ (absolute magnitude $>-19.5$ mag). It is only in the last dozen years that numerous examples of SNe that are substantially super-luminous ($>7\\times10^{43}$\\,erg\\,s$^{-1}$; $<-21$ mag absolute) were well-documented. Reviewing the accumulated evidence, we define three broad classes of super-luminous SN events (SLSNe). Hydrogen-rich events (SLSN-II) radiate photons diffusing out from thick hydrogen layers where they have been deposited by strong shocks, and often show signs of interaction with circumstellar material. SLSN-R, a rare class of hydrogen-poor events, are powered by very large amounts of radioactive $^{56}$Ni and arguably result from explosions of very massive stars due to the pair instability. A third, distinct group of hydrogen-poor events emits photons from rapidly-expanding hydrogen-poor material distributed over large radii, and are not powered by radioactivity (SLSN-I). These may be the hydrogen-poor analogs of SLSN-II. ",
        "introduction": "Supernova explosions play important roles in many aspects of astrophysics, being sources of heavy elements, ionizing radiation and energetic particles; driving gas outflows and shock waves that shape star and galaxy formation; and leaving behind compact neutron star and black hole remnants. The study of supernovae in general has thus been actively pursued for many decades.  The discovery of super-luminous supernova (SLSN; Figure~\\ref{LCfig}) events in the past decade is now focusing attention on these extreme explosions. The study of SLSNe is motivated, among other things, by their likely association with the deaths of the most massive stars; their potential contribution to the chemical evolution of the Universe and, at early times, to its reionization; and since they may be manifestations of physical explosion mechanisms that differ from those of their more common and less luminous cousins.  With extreme luminosities extending over tens of days (Fig.~\\ref{LCfig}) and, in some cases, copious ultra-violet (UV) flux, these events may become useful cosmic beacons to study distant star-forming galaxies and their gaseous environments; their long duration, coupled with the lack of long-lasting environmental effects, as well as the fact that they eventually disappear and allow their hosts to be studied without interference, offer some advantages over other probes of the distant universe such as short-lived gamma-ray burst (GRB) afterglows, and luminous high-redshift Quasars.  Accumulated observations suggest that SLSNe can be naturally grouped into three distinct subclasses. In analogy to lower-luminosity explosions, we split SLSNe into hydrogen-rich events (SLSN-II) and events lacking spectroscopic signatures of hydrogen (SLSN-I). The latter group is further divided into a minority of events whose luminosity is dominated by radioactivity (SLSN-R) and a distinct majority that require some other source of luminosity, which we simply refer to as SLSN-I. A detailed discussion of these three classes follows in section $\\S~3$ below.  Supernovae have been traditionally classified mainly according to their spectroscopic properties (see Filippenko 1997 for a review), while the luminosity of SNe does not play a role in the currently used scheme. In principle, almost all luminous SNe belong to one of two spectroscopic classes: Type IIn, hydrogen-rich events with narrow emission lines which are usually interpreted as signs of interaction with material lost by the star prior to the explosion; or Type Ic, events lacking hydrogen, helium, and strong silicon and sulphur lines around maximum, presumably associated with massive star explosions. However, as will be discussed below, the physical properties implied by the huge luminosities of SLSNe suggest they arise, in many cases, from progenitor stars that are very different from those of their much more common and less luminous analogs; we will propose a provisional extension of the classification scheme that can be applied to super-luminous events.  We consider below SNe with reported peak magnitudes M$<-21$\\,mag {\\it in any band} as being superluminous (Fig.~\\ref{LCfig}). See Supplementary Online Material (SOM) section $\\S~7$ for considerations related to determining this threshold.  As will be shown below, modern studies based on large SN samples and homogeneous, CCD-based luminosity measurements, show that SLSNe are very rare in nearby, luminous and metal-rich host galaxies (with the possible exception of luminous SNe IIn). One would therefore expect that older SN surveys, mainly conducted by monitoring nearby galaxies drawn from galaxy catalogs (dominated by luminous, metal-rich objects, we will refer to these as ``targeted surveys'' for short) would find very few luminous events. Indeed, Richardson et al. (2002) find no solid detections of super-luminous events resulting from the older targeted surveys (SOM $\\S~8$). A handful of events are determined to be ``genuinely overluminous'' (based on full, modern studies); Richardson et al. list SN 1997cy and SN 1999as; we will add SN 1999bd for which we present here new data. These events, found by the first generation of wide-field non-targeted surveys, are indeed the harbingers of the classes of SLSNe we can finally define today. We note that even the brightest SNe associated with cosmological Gamma-Ray Bursts (GRBs, e.g., the nickel-rich SN 1998bw) fall well below our SLSN threshold.  The structure of this review is as follows: we begin by a brief summary of the observational work leading to the discovery of superluminous SNe in section 2, and we synthesize the accumulated information and define classes of superluminous SNe in 3. We review open questions and controversies in 4, and conclude with a summary in 5.  \\begin{figure}[h] \\centering \\includegraphics[width=1\\textwidth]{LCfig.eps} \\caption{The luminosity evolution (light curve) of supernovae. Common SN explosions reach peak luminosities of $\\sim10^{43}$\\,erg\\,s$^{-1}$ (absolute magnitude $>-19.5$). The new class of super-luminous SN (SLSN) reach luminosities $\\sim10$ times higher. The prototypical events of the three SLSN classes (SLSN-I PTF09cnd, Quimby et al. 2011; SLSN-II SN 2006gy, Smith et al. 2007, Ofek et al. 2007, Agnoletto et al. 2009; and SN 2007bi, Gal-Yam et al. 2009) are compared with a normal Type Ia SN (Nugent template), Type IIn SN 2005cl (Kiewe et al. 2011), the average Type Ib/c light curve from Drout et al. (2012), the Type IIb SN 2011dh (Arcavi et al. 2011) and the prototypical Type II-P SN 1999em (Leonard et al. 2002). All data are in the observed $R$ band. See SOM for additional details.} \\label{LCfig} \\end{figure} ",
        "conclusions": "During the last dozen years, numerous super-luminous SN events have been discovered and studied. The accumulated data suggest these can be grouped into three distinct subclasses according to their observational and physical attributes. Radioactively-powered SLSN-R seem to be the best understood (and rarest) class, while hydrogen-rich SLSN-II and the most luminous hydrogen-poor SLSN-I are more common, but the physical origins of the extreme luminosity they emit is not clear at this time. With several ongoing surveys efficiently detecting additional examples, the amount of information about these objects, their rates and diversity, is likely to increase substantially in the coming few years."
    },
    "1208/1208.2027_arXiv.txt": {
        "abstract": "We study the abundance of satellite galaxies as a function of primary stellar mass using the SDSS/DR7 spectroscopic catalogue. In contrast with previous studies, which focussed mainly on bright primaries, our central galaxies span a wide range of stellar mass, $10^{7.5} \\leq M_*^{\\rm pri}/M_\\odot \\leq 10^{11}$, from dwarfs to central cluster galaxies. Our analysis confirms that the average number of satellites around bright primaries, when expressed in terms of satellite-to-primary stellar mass ratio ($m_*^{\\rm sat}/M_*^{\\rm pri}$), is a strong function of $M_*^{\\rm pri}$. On the other hand, satellite abundance is largely independent of primary mass for dwarf primaries ($M_*^{\\rm pri}<10^{10} \\, M_\\odot$). These results are consistent with galaxy formation models in the $\\Lambda$CDM scenario. We find excellent agreement between SDSS data and semi-analytic mock galaxy catalogues constructed from the Millennium-II Simulation.  Satellite galaxies trace dark matter substructure in $\\Lambda$CDM, so satellite abundance reflects the dependence on halo mass, $M_{200}$, of both substructure and galaxy stellar mass ($M_*$). Since dark matter substructure is almost scale-free, the dependence of satellite abundance on primary mass results solely from the well-defined characteristic mass in the galaxy mass-halo mass relation. On dwarf galaxy scales, where models predict a power-law scaling, $M_* \\propto M_{200}^{2.5}$, similarity is preserved and satellite abundance is independent of primary mass. For primaries brighter than the characteristic mass of the $M_*$-$M_{200}$ relation, satellite abundance increases strongly with primary mass. Our results provide strong support for the steep, approximately power-law dependence of dwarf galaxy mass on halo mass envisioned in $\\Lambda$CDM galaxy formation models. ",
        "introduction": "\\label{sec:intro}  Matching the galaxy luminosity function in the $\\Lambda$CDM scenario requires that the stellar mass of galaxies, $M_*$, should vary strongly with the virial\\footnote{Virial quantities are defined at the radius from the center of each halo where the mean enclosed density equals $200$ times the critical density of the Universe and are identified by a $200$ subscript. Units assume a Hubble constant of $H_0=73$ km s$^{-1}$ Mpc$^{-1}$ unless otherwise specified.} mass, $M_{200}$, of their surrounding dark matter halos. This exercise implies that the ``efficiency'' of galaxy formation, as measured by the ratio $M_*/M_{200}$, decreases steadily toward both smaller and larger masses from a maximum at $M_{200} \\sim 10^{12} M_\\odot$ \\citep{Yang2003,Vale2004,Shankar2006,Zheng2007,Kravtsov2010,Moster2010,Behroozi2010,Guo2010,Guo2011_sam}. On the scale of dwarf galaxies ($M_* < 10^{10}\\, M_\\odot$) these models require a near power-law dependence, $M_* \\propto M_{200}^{2.5}$, in order to reproduce observations of faint objects \\citep[however, see ][]{Behroozi2012}. Such a steep $M_*$-$M_{200}$ relation implies that dwarfs differing by as much as three decades in stellar mass (e.g., from the Fornax dwarf spheroidal to the Large Magellanic Cloud) inhabit halos spanning just over one decade in virial mass. Furthermore, very few galaxies exceeding $10^6 \\, M_\\odot$ are expected to have halos with virial masses below $10^{10}\\, M_\\odot$ \\citep{Ferrero2012}.  These predictions have been recently challenged by a series of observations, including (i) the lack of a characteristic velocity at the faint-mass end of blind HI surveys (expected if most dwarfs live in similar halos, \\citealt{Zwaan2010,Papastergis2011}); and (ii) the low virial mass (substantially below $10^{10}\\, M_\\odot$) inferred from dynamical data for the dwarf spheroidal companions of the Milky Way \\citep{Boylan-Kolchin2012} and for nearby dwarf irregulars \\citep{Ferrero2012}. The evidence, however, is indirect, since halo masses are estimated by extrapolating data that probe only the inner few kiloparsecs, where most baryons reside.  We explore here the possibility of using satellite galaxies to help constrain the virial masses of dwarf galaxies. The orbital motions of satellite companions have often been used to estimate halo masses \\citep[see, e.g.,][]{Zaritsky1997,Erickson1999,McKay2002,Prada2003,vandenBosch2004,Brainerd2005,Conroy2007,Wojtak2012}, but this work has largely been restricted to systems similar to or brighter than the Milky Way. This is partly due to the difficulties in obtaining redshifts for faint objects.  In addition, satellite companions are less common around dwarf galaxies than around larger systems: the number of satellites brighter than a certain fraction of the primary luminosity, $N(>L^{\\rm sat}/L^{\\rm pri})$, declines strongly toward fainter primaries \\citep[e.g. ][]{GuoQuan2011, Wang_White2012}. Dwarf galaxy associations do exist, but only a handful have been observed \\citep[e.g.,][]{Tully2006}.  In $\\Lambda$CDM, where satellite galaxies are thought to trace the substructure of cold dark matter halos, satellite systems are expected around all central galaxies, regardless of luminosity. The number of satellites, and their dependence on primary mass, should just reflect the abundance of substructure modulated by the dependence of galaxy formation efficiency on halo mass.  Substructure abundance has been studied extensively through numerical simulations, and shown to be nearly invariant with halo mass when expressed as a function of satellite mass normalized to that of the host \\citep{Moore1999b,kravtsov2004,Gao2004,Wang2012}. This result, together with the strict constraints on galaxy formation efficiency mentioned above, imply that satellite number counts provide useful information on the halo mass of dwarf galaxies. In particular, the near self-similarity of cold dark matter halos provides an instructive test: if satellite galaxies trace substructure then the abundance of {\\it luminous} satellites should also be scale-free on scales where galaxy mass and halo mass are related by a power law.   We explore these issues here by identifying primary-satellite systems in galaxy catalogues constructed from the Sloan Digital Sky Survey and by comparing them with predictions from a semianalytic mock galaxy catalogue based on the Millennium Simulations. The plan for this paper is as follows. Sec.~\\ref{sec:data} describes briefly the observational and model datasets while Sec.~\\ref{sec:results} presents our main results. We summarize our main conclusions in Sec.~\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We study the abundance of satellites as a function of primary stellar mass in the Sloan Digital Sky Survey. Using the SDSS/DR7 spectroscopic sample from the NYU-VAGC catalogue we are able to extend previous studies to significantly fainter primaries, $M_*^{\\rm pri}=[10^{7.5}$-$10^{11}] M_\\odot$.  In agreement with previous work, we find that the abundance of satellites exceeding a given satellite-to-primary stellar mass ratio, $m_*^{\\rm sat} /M_*^{\\rm pri}$, depends strongly on $M_*^{\\rm pri}$ for bright primaries.  On the other hand, {\\it the abundance of satellites around dwarf primaries, $M_*^{\\rm pri} < 10^{10} M_\\odot$, is approximately independent of primary stellar mass}.  These results are in excellent agreement with predictions of semi-analytic models within $\\Lambda$CDM. These trends arise from the mass invariance of substructure in CDM halos and from the varying efficiency of galaxy formation as a function of halo mass. On dwarf galaxy scales, where the relation between galaxy mass and halo mass is well approximated by a steep power law, the invariance of satellite abundance with primary mass reflects directly the scale-free nature of substructure. Around bright galaxies the scaling between galaxy mass and halo mass deviates from a simple power law, leading to the observed strong increase of satellite abundance with increasing primary mass.   Some caveats on the statements above need to be mentioned.  The first relates to the definition of halo mass for a galaxy that has entered the virial radius of a larger object.  The mass in subhalos is in general ill-defined, depending not only on the identification algorithm \\citep{Knebe2011}, but also on time, due to tidal stripping \\citep[e.g. ][]{Tormen1998, Klypin1999b, Hayashi2003}. The stellar mass, however, is more resilient to tidal effects and remains approximately constant after accretion onto the host \\citep{White_Rees1978, Sales2007a, Penarrubia2008}.  Thus for satellite galaxies, the virial mass at the moment of accretion is more closely related to the stellar mass than their present-day dark halo mass.  Since the abundance of subhaloes according to their infall mass --termed the ``unevolved'' subhalo mass function-- is also independent of host halo mass when written as a function of the relative mass between satellite and host $m_{\\rm acc}^{\\rm sub}/M_{\\rm DM}^{\\rm host}$ \\citep[e.g. ][]{Giocoli2008, Yang2011}, the arguments given above remain valid.   The second caveat involves the possible dependence of the $M_*$-$M_{200}$ relation on redshift and, related to that, whether satellites and primaries follow the same relation between stellar mass and halo mass.  Abundance matching models suggest that the link between the stellar mass and halo mass evolves weakly with redshift \\citep{Shankar2006,Conroy2009,Wang_Jing2010,Leauthaud2011,Yang2012,Moster2012,Behroozi2012}. Given that surviving satellites are preferentially accreted at late times \\citep[e.g. ][]{Gao2004,Sales2007a} the $M_*$-$M_{200}$ relation for satellites and centrals are expected to be similar. In particular, self-consistent semi-analytical modelling of galaxies shows only small differences between the two \\citep[e.g. see Fig. 9 of ][]{Guo2011_sam}.  Notice that these arguments do not allow for stripping of the stars from satellites. Arguably, this could complicate the evolution. We note, however, that numerical models for dwarf galaxies suggest that once stellar stripping sets in, total disruption soon follows \\citep{Penarrubia2008}.  We therefore expect partial stripping of stars to have minor effects in large statistical samples. On the other hand, several studies have suggested that accounting for total disruption of satellites is needed to match observations \\citep[e.g. ][]{Weinmann2006, Kimm2009}.  We address this by comparing observational results with a semi-analytic model that explicitly treats satellite disruption by tidal forces.   Under the assumption of a $\\Lambda$CDM universe, the good agreement in shape and normalization between satellite counts in SDSS and those in the mock catalogue provides support for a steep stellar-halo mass relation for dwarfs, consistent with the $M_* \\propto M_{200}^{2.5}$ predicted both by semi-analytic models and by extrapolations of current abundance-matching analyses.  More definitive constraints on the slope of the $M_*$-$M_{200}$ relation for dwarf galaxies may come from a robust determination of the slope of satellite abundances around isolated primaries in tandem with studies of the effect of scatter in the stellar mass - halo mass relation. Probing increasingly fainter companions in observational surveys of the surroundings of isolated dwarfs may prove crucial for this goal."
    },
    "1208/1208.0352_arXiv.txt": {
        "abstract": "We report the identification of the M9 dwarf SDSS~J000649.16$-$085246.3 as a spectral binary and radial velocity variable with components straddling the hydrogen burning mass limit. Low-resolution near-infrared spectroscopy reveals spectral features indicative of a T dwarf companion, and spectral template fitting yields component types of M8.5$\\pm$0.5 and T5$\\pm$1. High-resolution near-infrared spectroscopy with Keck/NIRSPEC reveals pronounced radial velocity variations with a semi-amplitude of 8.2$\\pm$0.4~{\\kms}.  From these we determine an orbital period of 147.6$\\pm$1.5~days and eccentricity of 0.10$\\pm$0.07, making SDSS~J0006$-$0852AB the third tightest very low mass binary known. This system is also found to have a common proper motion companion, the inactive M7 dwarf LP~704-48, at a projected separation of 820$\\pm$120~AU. The lack of H$\\alpha$ emission in both M dwarf components indicates that this system is relatively old, as confirmed by evolutionary model analysis of the tight binary. LP~704-48/SDSS~J0006$-$0852AB is the lowest-mass confirmed triple identified to date, and one of only seven candidate and confirmed triples with total masses below 0.3~{\\msun} currently known. We show that current star and brown dwarf formation models cannot produce triple systems like LP~704-48/SDSS~J0006$-$0852AB, and we rule out Kozai-Lidov perturbations and tidal circularization as a viable mechanism to shrink the inner orbit. The similarities between this system and the recently uncovered low-mass eclipsing triples NLTT~41135AB/41136 and LHS~6343ABC suggest that substellar tertiaries may be common in wide M dwarf pairs. ",
        "introduction": "Our current theoretical and observational understanding of star formation holds that most stars form in multiple systems, with the frequency, mass ratio distribution and period distribution of multiples varying as a function of mass and possibly formation environment (e.g., \\citealt{2004MNRAS.351..617D,2007ApJ...671.2074A,2007prpl.conf..427B,2009MNRAS.392..590B,2012MNRAS.419.3115B,2011ApJ...731....8K}).  Dynamics play an important role in multiple formation, as both the fragmentation of collapsing cloud cores \\citep{2001ApJ...551L.167B} and massive circumstellar disks \\citep{2009MNRAS.392..413S,2011MNRAS.413.1787S} give rise to small N groups that either dissolve or stabilize into hierarchical systems \\citep{2003A&A...400.1031S}.  Numerical calculations show that both fragmentation and dynamical evolution of stellar multiples are sensitive to the initial cloud properties---geometry, density, turbulence spectrum, and radiative and external feedback---so multiplicity statistics provide important tests for star formation theory, on par with the initial mass function, stellar mass segregation and velocity distributions.  The mass dependence of multiplicity among stars is well established, with observed frequencies ranging from $>$80\\% for OBA dwarfs \\citep{2002A&A...382...92S,2005A&A...430..137K,2007ApJ...670..747K} to $\\sim$30\\% for systems with M dwarf primaries \\citep{1992ApJ...396..178F,1997AJ....113.2246R,2004ASPC..318..166D}.  The frequency of multiples appears to drop even further in the very low mass (VLM; M $\\lesssim$ 0.1~{\\msun}) stellar and substellar regimes.  Resolved imaging studies of field and cluster VLM dwarfs, based on ground-based adaptive optics (AO) and Hubble Space Telescope (HST) observations, find relatively consistent (volume-limited) binary frequencies of 10--20\\% (e.g., \\citealt{2003AJ....126.1526B,2003ApJ...586..512B,2006ApJS..166..585B,2003ApJ...587..407C}).  However, estimates for the frequency of tightly-bound spectroscopic VLM binaries (roughly 10\\% of currently known systems) span a broad range of 1--25\\% \\citep{2005MNRAS.362L..45M,2006AJ....132..663B,2006MNRAS.372.1879K,2008A&A...492..545J,2010ApJ...723..684B,2012ApJ...744..119C}.  As such, there remains significant uncertainty in the overall VLM binary frequency, which has a direct impact on our understanding of how brown dwarfs form in the first place (e.g., \\citealt{2003MNRAS.342..926D}). There is also evidence that the orbital characteristics of multiples are mass-dependent, with VLM binaries being on average tighter ($\\langle{a}\\rangle \\approx$ 7~AU versus $\\approx$ 30~AU) and more frequently composed of equal-mass components ($\\langle{q}\\rangle \\equiv {\\rm M}_2/{\\rm M}_1 \\approx 1$ versus $\\approx$ 0.3) than their solar-mass counterparts \\citep{2007ApJ...668..492A}.  Again, these trends may be skewed by biases inherent to the known sample of VLM binaries, identified largely as imaged pairs.  A complete understanding of VLM multiplicity statistics therefore requires more robust constraints on the frequency and characteristics of short-period and low-$q$ multiples \\citep{2007prpl.conf..427B}.  An alternative method for identifying VLM multiples that avoids separation biases (both intrinsic and projected) is through spectral binaries, systems with components of different spectral types whose combined-light spectra exhibit distinct peculiarities.  While this technique is more commonly associated with white dwarf-M dwarf pairs (e.g., \\citealt{2006AJ....131.1674S}), spectral binaries containing late M or L dwarf primaries and T dwarf secondaries have also been recognized due to their unique and highly structured near-infrared spectra  \\citep{2004ApJ...604L..61C,2007AJ....134.1330B,2008ApJ...681..579B,2008arXiv0811.0556S,2010ApJ...710.1142B,2011ApJ...739...49B,2010AJ....140..110G,2011ApJ...732...56G}.  Because blended light systems can be identified at any separation, M/L+T spectral binaries can probe the very closest separations that are inaccessible to direct imaging studies ($a \\lesssim$ 50~mas). An illustrative case is the M8.5 + T5 binary 2MASS~J03202839$-$0446358 (hereafter 2MASS~J0320$-$0446), a system independently identified as both a spectral binary \\citep{2008ApJ...681..579B} and a radial velocity (RV) variable with a period of 0.68~yr and primary semi-major axis $a_1\\sin{i}$ = 0.157$\\pm$0.003~AU  \\citep{2008ApJ...678L.125B,2010ApJ...723..684B}. With a projected separation of $\\approx$17~mas, this system is unresolvable with current imaging and interferometric technology. Moreover, the significant difference in component magnitudes ($\\Delta{K}$ = 4.3$\\pm$0.6) makes this system a potential low $q$ pair. The independent constraints provided by the component classifications and RV orbit yield robust limits on the age ($\\gtrsim$2~Gyr) and orbital inclination ($\\gtrsim 53\\degr$) of this system \\citep{2009AJ....137.4621B}.  In an effort to confirm and characterize spectral binaries identified in low-resolution near-infrared spectroscopy, we have identified a new VLM system exhibiting significant RV variability.  The source, {\\name} (hereafter {\\namesh}; \\citealt{2008AJ....135..785W}) is an optically-classified M9 which appears to be both a short-period (0.4~yr) binary with M8.5 and T5$\\pm$1 components, and a co-moving wide companion to the inactive M7 dwarf LP~704-48.  Together, this system comprises the lowest-mass triple confirmed to date. In Section~2 we describe our optical and near-infrared imaging and spectroscopic observations of this system using the 1m Lick Observatory Nickel Direct Imaging Camera, the 3m NASA Infrared Telescope Facility (IRTF) SpeX spectrograph \\citep{2003PASP..115..362R}, the 4m Kitt Peak National Observatory (KPNO) RC spectrograph, and the 10m Keck II NIRSPEC spectrograph \\citep{1998SPIE.3354..566M}. In Section~3 we examine the properties of the {\\namesh} and LP~704-48 pair, assessing their common proper motion, distance and probability of chance alignment; as well as the age and metallicity of both components as derived from spectroscopic and kinematic indicators. In Section~4 we analyze the low-resolution near-infrared spectrum of {\\namesh} to infer the presence of, and characterize, its T dwarf companion. In Section~5 we analyze our RV measurements which allow us to extract both the orbital properties (including constraints on the orbital inclination) and verify the relatively old age for this system. In Section~6 we discuss LP~704-48/{\\namesh}AB in the context of other hierarchical low-mass triples, and examine whether current star formation theories or three-body interactions can create such systems. Our results are summarized in Section~7. ",
        "conclusions": "\\subsection{Very Low Mass Triples}  The LP~704-48/{\\namesh}AB system joins a growing list of VLM triples in which all three components have masses near or below 0.1~{\\msun} (assuming a mass of 0.093~{\\msun} for LP~704-48 based on the evolutionary models used above). Formally defining the category ``VLM triple'' as a bound system of three hydrogen-rich objects with a total mass of 0.3~{\\msun} or less, we list the seven candidate and confirmed systems that fall into this category in Table~\\ref{tab:triples}. With the exception of LHS~1070ABC, these systems are strongly hierarchical, with a ratio of outer to inner separations ranging from $\\sim$5 for LP~714-37ABC to $\\sim$2900 for LP~704-48/{\\namesh}AB. They also have estimated component masses that are comparable to each other, although this may be a selection bias (sensitivity limits) and subject to the typically large uncertainties in mass estimates for substellar dwarfs. Three of these systems have total estimated masses M$_{tot}$ $\\lesssim$ 0.15~{\\msun}, although their tertiary components have not been solidly confirmed. DENIS-P~J020529.0-115925ABC (hereafter DENIS~J0205-1159AB(C); \\citealt{1997A&A...327L..25D,1999ApJ...526L..25K,2005AJ....129..511B}) was resolved into a $\\sim$7~AU binary by AO and HST imaging. Evidence for a third component was reported by \\citet{2005AJ....129..511B}, but only on the basis of persistent residuals in point spread function fitting of the resolved pair in multi-epoch HST images. 2MASS J08503593+1057156ABC (hereafter 2MASS~J0850+1057AB(C); \\citealt{1999ApJ...519..802K,2001AJ....121..489R}) was resolved as a 0$\\farcs$16 binary with HST, and subsequently identified as a candidate triple based on the later classification of its primary, despite this component being roughly a magnitude brighter in the near-infrared \\citep{2011AJ....141...70B}. This result has been called into question by \\citet{2012arXiv1201.2465D}, however, based on reanalysis of the combined-light spectrum. Kelu~1AB(C) \\citep{1997ApJ...491L.107R,1999Sci...283.1718M,2005ApJ...634..616L,2006PASP..118..611G} was resolved as a 6~AU visual double with AO and HST imaging, and its  primary identified as an unresolved ($<$4~AU) L/T spectral binary by \\citet{2008arXiv0811.0556S}, although there has been no evidence of RV variability in the combined-light system at the 1--3~{\\kms} level \\citep{2006AJ....132..663B,2010ApJ...723..684B}. Given the uncertain nature of these triple candidates, we argue that LP~704-48/{\\namesh}AB is currently the lowest-mass triple verified through multiple techniques.  It is also the only VLM triple system in Table~\\ref{tab:triples} for which multiple orbits of the tight inner binary have been observed.  \\begin{deluxetable*}{lllllccccccl} \\tablecaption{Confirmed and Candidate Very Low Mass Triples (M$_{tot} \\lesssim 0.3$~{\\msun}) \\label{tab:triples}} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablehead{ \\multicolumn{2}{c}{Components\\tablenotemark{a}}  & \\multicolumn{3}{c}{Spectral Types}  & \\multicolumn{4}{c}{Estimated Masses ({\\msun})} & \\multicolumn{2}{c}{Separations (AU)}  & \\colhead{Ref} \\\\ \\colhead{A} & \\colhead{BC} & \\colhead{A} & \\colhead{B} & \\colhead{C} & \\colhead{A} & \\colhead{B} & \\colhead{C} &  \\colhead{Total} & \\colhead{A-BC} & \\colhead{BC} & \\\\ } \\startdata LP~714-37A & LP~714-37BC & M5.5 & M8 & M8.5 & 0.11 & 0.09 & 0.08 & 0.28 & 36$\\pm$5 & 6.8$\\pm$0.9 & 1 \\\\ LHS~1070A & LHS~1070BC & M5.5 & M8.5 & M9 & 0.115  &  0.08 & 0.077 & 0.27 & $\\sim$12\\tablenotemark{b} &  3.57$\\pm$0.07 & 2,3 \\\\ LP~213-68 & LP~213-67AB & M8 & M8 & L0 & 0.09  & 0.09 & 0.084 & 0.27 & 340$\\pm$60 & 2.9$\\pm$0.6 & 4,5 \\\\ {\\bf LP~704-48} & {\\bf SDSS~J0006$-$0852AB} & M7 & M8.5 & T5 & 0.092 & 0.083 & 0.056 & 0.23 & 820$\\pm$120 & 0.286$\\pm$0.009 & 6 \\\\ DENIS~J0205-1159A & DENIS~J0205-1159B(C)\\tablenotemark{c} &  L5 & L8 & T0: & 0.06 & 0.05 & 0.04 & 0.15 & 7$\\pm$1 & $\\sim$1.3 & 7 \\\\ 2MASS~J0850+1057B & 2MASS~J0850+1057A(C)\\tablenotemark{c} & L6 & L7 & L7: & 0.05 & 0.05 & 0.05 & 0.15 &  6.0$\\pm$0.9 & $<$4  & 8 \\\\ Kelu~1B & Kelu~1A(C)\\tablenotemark{c} & L3p & L0.5 & T7: & 0.05 & 0.06 & 0.02 & 0.13 &  6.4$^{+2.4}_{-1.3}$ & $<$4  & 9 \\\\ \\enddata \\tablenotetext{a}{For this table, we refer to the outermost component as ``A'' and the inner binary as ``BC'' irrespective of mass or designation.} \\tablenotetext{b}{Scaling from the inner semi-major axis using the period-mass ratio: $(a_{out}/a_{in})^3 = (P_{out}/P_{in})^2 \\times (M_{out}/M_{in})$, and values from Seifahrt et al.~(2008).} \\tablenotetext{c}{Candidate triples with unconfirmed third component.} \\tablerefs{ (1) \\citet{2006ApJ...645L.153P}; (2) \\citet{2001A&A...367..183L}; (3) \\citet{2008A&A...484..429S}; (4) \\citet{2000MNRAS.311..385G}; (5) \\citet{2003ApJ...587..407C}; (6) This paper; (7) \\citet{2005AJ....129..511B}; (8) \\citet{2011AJ....141...70B}; (9) \\citet{2008arXiv0811.0556S}.} \\end{deluxetable*}  The LP~704-48/{\\namesh}AB system shares much in common with two recently uncovered, but slightly more massive, low-mass triples containing transiting substellar components: NLTT~41135AB/NLTT~41136 \\citep{2010ApJ...718.1353I} and LHS~6343ABC \\citep{2011ApJ...730...79J}. Both are similarly composed of relatively wide M-dwarf pairs (55~AU and 20~AU, respectively) with one component hosting a tightly-orbiting (0.02~AU and 0.08~AU) substellar mass (0.03~{\\msun} and 0.06~{\\msun}) companion. For these systems, the component separations are roughly an order of magnitude smaller than for LP~704-48/{\\namesh}AB, but the relative inner to outer separations and component masses are comparable.  Moreover, based on evolutionary models, the tertiaries of these systems are also likely to be T dwarfs. The fact that three such systems have been identified in the span of three years suggests that such low-mass triple configurations may be quite common, particularly among wide VLM pairs (e.g., \\citealt{2010ApJ...720.1727L}), although a robust survey is needed to quantify the incidence of such systems.   \\subsection{Stability of the LP~704-48/{\\namesh}AB System}  The LP~704-48/{\\namesh}AB system is at an extremum among the VLM triples listed in Table~\\ref{tab:triples} in that it has both the widest outer separation and the smallest (measured) inner separation in the sample.  The outer pairing is remarkable for a system with M$_{tot}$ $\\lesssim$0.25~{\\msun}; currently, only three other VLM field binaries are known to have projected separations $>$500~AU.\\footnote{Koenigstuhl-1 at 1800~AU \\citep{2007A&A...462L..61C}, 2MASS~J0126555$-$502239 at 5100~AU; \\citep{2007ApJ...659L..49A,2009ApJ...692..149A} and 2MASS~J12583501+4013083 at 6700~AU \\citep{2009ApJ...698..405R}.} The gravitational binding energy of LP~704-48 and {\\namesh}AB, $|E_b| \\lesssim$ (2--3)$\\times$10$^{41}$~erg, is low but not unprecedented.  Recently, several wide multiples of comparable total mass and binding energy have been found (e.g., \\citealt{2010AJ....139.2566D,2010ApJ...720.1727L}), including high-order systems such as NLTT~20346AB/2MASS~J0850359+105716AB(C) ($|E_b| \\approx 0.4\\times10^{41}$~erg; \\citealt{1999ApJ...519..802K,2001AJ....121..489R,2011AJ....141...70B,2011AJ....141...71F}) and G~124-62/DENIS-P~J144137.3$-$094559AB ($|E_b| \\approx 3\\times10^{41}$~erg; \\citealt{1999AJ....118.2466M,2006A&A...456..253M,2003AJ....126.1526B,2005A&A...440..967S}), both of which contain substellar components. In part, it is the additional mass of the T dwarf tertiary that pushes the LP~704-48/{\\namesh}AB system into a ``normal'' regime in mass/separation space (see \\citealt{2011AJ....141...71F}). This component also contributes to the long-term stability of the wide pair to external perturbation \\citep{1987ApJ...312..367W,2010AJ....139.2566D}.  \\begin{figure} \\centering \\epsscale{1.1} \\plotone{f14.eps} \\caption{Stability ratio $Y/Y_{min}$ as a function of total system mass for a sample of low mass triple systems drawn from the Multiple Star Catalog \\citep{1997A&AS..124...75T}; \\citet{2011AJ....141...71F}; and this paper.  The LP~704-48/{\\namesh}AB system is indicated by the large red circle at left; the eclipsing triples NLTT~41135AB + NLTT~41136 \\citep{2010ApJ...718.1353I} and LHS~6343ABC \\citep{2011ApJ...730...79J} are indicated by blue squares. \\label{fig:hierarchical}} \\end{figure}  Because of the large difference in outer to inner separations, LP~704-48/{\\namesh}AB is also exceptionally stable to internal dynamical disruption. Internal stability can be quantified with the criteria of \\citet{1995ApJ...455..640E}, who examined the ratio of outer periapse $a_{out}(1-e_{out})$ to inner apoapse $a_{in}(1+e_{in})$ for triple star systems.  For simplicity, we drop the eccentricity terms and approximate the ratio $Y = a_{out}/a_{in}$, using as $a_{out}$ the projected separation of the wide pairing.  For LP~704-48/{\\namesh}AB, $Y \\approx 2900$, a value several orders of magnitude greater the critical ratio \\begin{equation} Y_{min} = 1+\\frac{3.7}{q_{out}^{1/3.}} + \\frac{2.2}{1+q_{out}^{1/3}} + 1.4q_{in}^{1/3}\\frac{q_{out}^{1/3} - 1}{1+q_{out}^{1/3}}, \\end{equation} which is $\\approx$5 for this system, assuming $q_{in} \\equiv {\\rm M}_B/{\\rm M}_A \\approx 0.7$ for {\\namesh}AB and $q_{out} \\equiv ({\\rm M}_A+{\\rm M}_B)/{\\rm M}_{LP~704-48} \\approx 1.6.$\\footnote{These estimates also imply a period ratio $X = Y^{1.5}\\left(\\frac{q_{out}}{q_{out}+1}\\right)^{0.5} \\approx 10^5$.} Figure~\\ref{fig:hierarchical} compares the ratio $Y/Y_{min}$ for LP~704-48/{\\namesh}AB and other low-mass triples listed in the most recent version (April 2010) of the Multiple Star Catalog \\citep{1997A&AS..124...75T}, \\citet{2011AJ....141...71F} and Table~\\ref{tab:triples}.  By this metric, LP~704-48/{\\namesh}AB stands out as the most internally stable low-mass triple known, slightly more stable than the NLTT~41135AB/NLTT~41136 system \\citep{2010ApJ...718.1353I}.  With such a large semi-major axis ratio, we note that N-body simulations find the outer orbits of such systems are likely to have a significant eccentricity \\citep{2002A&A...384.1030S}, so the ratio of outer periapse to inner apoapse may be smaller than the semimajor axis ratio used here.  \\subsection{On the Formation of LP~704-48/{\\namesh}AB}  Given the now established existence of a handful of VLM triples like LP~704-48/{\\namesh}AB, it is worth examining whether such systems and their characteristics match the predictions of star and brown dwarf formation models. We emphasize that current samples are far from statistically robust or complete, and detection biases may be significant.  Despite the low production rate of VLM multiples in general, current models do produce VLM higher-order systems.  A commonly-ascribed mechanism is dynamical scattering in the post-accretion phase of a young cluster \\citep{2003A&A...400.1031S,2004MNRAS.351..617D,2004A&A...414..633G}. In their study of small N-body cluster dynamics, \\citet{2003A&A...400.1031S} found that 11\\% of their simulated systems formed triple and higher-order multiples, and over 40\\% of the triples contained at least one brown dwarf component.  Most of these systems were widely separated and hierarchical, typically with the distant companion being of low mass (in systems containing two brown dwarfs, the distant companion was a brown dwarf in 90\\% of cases).   However, all wide triples ($a \\gtrsim$ 100~AU) with brown dwarf components had extremely low mass ratios ($q < 0.3$), whereas the systems listed in Table~\\ref{tab:triples} are largely composed of near-equal mass components.  Moreover, this simulation proved unable to generate VLM binaries tighter than $\\sim$1~AU, and the authors conclude that dynamical interactions alone cannot explain the existence of such systems (which at the time included only PPl~15).  Hence, it would appear that dynamics alone is incapable of producing VLM multiples like LP~704-48/{\\namesh}AB, a conclusion that has been reached for more massive higher-order multiples as well \\citep{2008MNRAS.389..925T}.  Accretion plays an important role in the formation of multiple systems, as a mechanism for angular momentum exchange and dissipation, and driving binaries to near-equal masses (e.g., \\citealt{2002MNRAS.336..705B}). \\citet{2012MNRAS.419.3115B} recently examined VLM multiplicity in radiative hydrodynamic simulations of molecular cloud collapse, which incorporates both gas accretion (with some radiative feedback) and dynamics, albeit over shorter periods than pure N-body simulations due to computational constraints. The Bate study produced three low-mass triple systems (M$_{tot}$ $\\lesssim$ 0.6~AU), all of which were hierarchical with either the high mass or middle mass component at wide separation. While an exact clone to LP~704-48/{\\namesh}AB was not generated, one system composed of a 0.15~{\\msun} plus 0.03~{\\msun}, 14~AU inner binary with a 0.07~{\\msun} wide companion at 194~AU is a reasonable analog, although the large inner eccentricity of this triple ($e$ = 0.8) may ultimately lead to its dissolution.  One fundamental shortcoming of this study is the $\\sim$1~AU separation range limit on close interactions, again due to computational constraints. \\citet{2009MNRAS.392..590B} examined VLM binary statistics at ten times higher resolution but over a shorter timeframe and employing slightly different gas physics. While they could successfully form 0.1--1.0~AU VLM binaries in these simulations, no analysis was made of higher-order multiples.  The fragmentation of circumstellar disks around massive stars is another proposed mechanism for forming VLM stars and brown dwarfs \\citep{2008A&A...480..879S,2009MNRAS.392..413S}.  Unfortunately, the predicted binary rate for VLM systems created in these environments is very low ($\\sim$8\\% for ejected systems) and no low-mass triples have been created in simulations to date.  Furthermore, while the separation distribution of simulated substellar pairs peaks in a range that is close to that inferred for {\\namesh}AB (0.3 $< a <$ 0.6~AU), these binaries tend to have very large eccentricities ($e > 0.7$), a feature not seen in any of  the VLM spectroscopic binaries identified to date.  As such, disk fragmentation does not appear to be a viable mechanism for making triples like LP~704-48/{\\namesh}AB.  It is important to emphasize that current models fall short in reproducing very short-period, low-eccentricity VLM pairs like {\\namesh}AB and 2MASS~J0320$-$0446AB. Such systems are necessarily the product of dynamical and dissipative evolution, as opacity-limited fragmentation constrains the initial separations of self-gravitating masses to $\\gtrsim$10~AU \\citep{1969MNRAS.145..271L,1976MNRAS.176..367L}. \\citet{2002MNRAS.336..705B} resolves this problem through a combination of dynamical interactions and accretion from a circumbinary/circumtertiary disk, although again these simulations fail to produce stable VLM multiples. Indeed, three-body encounters more often than not disrupt VLM multiples \\citep{2010MNRAS.404..721M}. A more intriguing mechanism for {\\namesh}AB is Kozai-Lidov eccentricity perturbations induced by LP~704-48, followed by circularization through tidal friction (KCTF; \\citealt{1962AJ.....67..591K,1968AJ.....73..190H,1998MNRAS.300..292K,2007ApJ...669.1298F,2012ApJ...750..106S}).  This mechanism has been proposed to explain the high fraction of spectroscopic binaries with tertiary companions (96\\% for SBs with $P < $3~dy; \\citealt{2006A&A...450..681T}; see also \\citealt{2010ApJS..190....1R}).  KCTF predicts end states with roughly circular inner orbits and large period ratios ($X \\gtrsim 10^3-10^4$; \\citealt{2007ApJ...669.1298F}), similar to {\\namesh}AB.  Unfortunately, the timescale for eccentricity pumping is long for LP~704-48/{\\namesh}AB ($\\tau {\\sim} P^2_{out}/P_{in}$ $\\approx$ 3~Gyr), while tidal circularization is extremely inefficient at the current orbital separation of {\\namesh}AB\\footnote{Using Eqns~50 and 52 from \\citet{1981A&A....99..126H}, the timescale for circularization: \\begin{equation} t_{circ} = \\frac{0.0024}{kq}\\left(\\frac{a_0}{R_1}\\right)^8\\frac{P_0}{\\tau}P_0 \\end{equation} where $k$ is the primary's Love number, $q$ the mass ratio, $a_0$ and $P_0$ the final semimajor axis and period, $R_1$ the primary's radius and $\\tau$ the tidal lag with respect to the displacement between primary and secondary.  Assuming $k \\sim 0.5$ (i.e., comparable to Jupiter and Saturn), $q = 0.7$, $a_0/R_1$ = 0.29~AU/$R_{Jupiter} \\approx 620$ and $P_0$ = 0.4~yr and $\\tau \\approx 0.1P_0$, we derive $t_{circ} \\approx 6{\\times}10^{20}$~yr, far too long to have played a role in the dynamic evolution of {\\namesh}AB.}  \\citep{1981A&A....99..126H,2005ApJ...620..970M}.  Thus, it appears that KCTF is probably not responsible for the current configuration of this system, nor presumably for the isolated tight VLM binaries PPl~15AB, 2MASS~J0535$-$0546AB or 2MASS~J0320$-$0446AB.  A final possibility is that LP~704-48 has played no role in the evolution of the {\\namesh}AB binary, which could have shrunk and circularized through dissipative interactions with an early circumbinary disk or a closer encounter resulting in an ejection.  The wider pairing with LP~704-48 may simply be a normal (albeit rare) outcome of cloud fragmentation or parallel trajectories during the dissolution of the natal cluster \\citep{2010MNRAS.404.1835K,2010MNRAS.404..721M}. LP~704-48 and {\\namesh}AB may be more akin to stellar cousins that stellar siblings."
    },
    "1208/1208.2999_arXiv.txt": {
        "abstract": "We report the discovery of 34 stars in the Hamburg/ESO Survey for metal-poor stars and the Sloan Digital Sky Survey that have [Fe/H] $\\la$ --3.0.  Their median and minimum abundances are [Fe/H] = --3.1 and --4.1, respectively, while 10 stars have [Fe/H] $<$ --3.5. High-resolution, high-$S/N$ spectroscopic data -- equivalent widths and radial velocities -- are presented for these stars, together with an additional four objects previously reported or currently being investigated elsewhere.  We have determined the atmospheric parameters, effective temperature (\\teff) and surface gravity (\\logg), which are critical in the determination of the chemical abundances and the evolutionary status of these stars.  Three techniques were used to derive these parameters.  Spectrophotometric fits to model atmosphere fluxes were used to derive {\\teff}, {\\logg}, and an estimate of $E(B-V)$; H$\\alpha$, H$\\beta$, and H$\\gamma$ profile fitting to model atmosphere results provided the second determination of {\\teff} and {\\logg}; and finally, we used an empirical \\teff\\--calibrated H$\\delta$ index, for the third, independent \\teff\\ determination. The three values of \\teff\\ are in good agreement, although the profile fitting may yield systematically cooler \\teff\\ values, by $\\sim100$K.  This collective data set will be analyzed in future papers in the present series to utilize the most metal-poor stars as probes of conditions in the early Universe. ",
        "introduction": "\\label{Sec:intro}  Six decades after the discovery of metal-poor stars by \\citet{chamberlain51}, the study of these objects has become a mature area of research.  In 2010, high-resolution, high-signal-to-noise ($S/N$) chemical abundance analyses existed for some 400 stars that have [Fe/H] $< -2.5$, some 24 with [Fe/H] $< -3.5$, and three having [Fe/H] $< - 4.5$ (see the compilation of \\citet{frebel10}\\footnote{https://www.cfa.harvard.edu/$\\sim$afrebel/}) . As discussed by many authors, the most metal-poor stars, believed to have formed at redshifts z $\\ga$~6, are among the best probes of conditions in the early Universe, including in particular the formation of the first stars and the first chemical elements.  We shall not repeat here the case for this endeavor, but refer the reader to earlier works (see, e.g., \\citealt{bessell84, mcwilliam95, ryanetal96, norris01, cayrel04, beers05, cohen08, lai08, frebel11}) for a thorough discussion of the rationale that drives this very active field.  The aim of the present series of papers is to increase the inventory of the most metal-poor stars, by which we mean [Fe/H] $\\la -3.5$, in order to understand the Metallicity Distribution Function (MDF) at lowest metallicity and to investigate abundance patterns that contain clues to conditions at the earliest times.  We are interested to gain insight into the relative abundances ([X/Fe]\\footnote{[X/Fe] = $\\log_{10}(\\rm N_{\\rm X}/N_{\\rm Fe})_\\star - \\log_{10}(N_{\\rm X}/N_{\\rm Fe})_\\odot$}) in this abundance regime, involving in particular C, N, O, Mg, Na, Al, and the heavy neutron-capture elements, which deviate strongly from the relatively well-defined behavior of the majority of stars in the ([X/Fe], [Fe/H])-- planes at higher abundances ([Fe/H] $>$ --3.5) (see, e.g., \\citealt[Figure~11]{norris07}).  Section 2 of the paper is concerned with the discovery of stars having [Fe/H] $< -3.0$, by obtaining medium-resolution spectra of candidate metal-poor stars.  Section 3 describes follow-up high-resolution, high-$S/N$ spectroscopy of the most metal-poor of them, together with the reduction and analysis of the spectra to produce the basic data (equivalent widths and radial velocities) upon which subsequent analyses will be based.  Our sample contains 38 stars, some 13 of which (12 identified in the endeavors described in Section 2) have [Fe/H] $< -3.5$, based on subsequent high-resolution, high-$S/N$ spectroscopic analysis.  In Section IV we describe the determination of accurate effective temperatures needed for the chemical abundance analysis.  Papers II, III (Yong et al.\\, 2012a, b) and Papers IV, V (Norris et al.\\, 2012, b) will present accurate chemical abundances and discussion of the elements from lithium through to the heavy-neutron-capture elements.  \\newpage ",
        "conclusions": ""
    },
    "1208/1208.0164_arXiv.txt": {
        "abstract": "We discuss the design and expected performance of STRIP (STRatospheric Italian Polarimeter), an array of coherent receivers designed to fly on board the LSPE (Large Scale Polarization Explorer) balloon experiment. The STRIP focal plane array comprises 49 elements in Q band and 7 elements in W-band using cryogenic HEMT low noise amplifiers and high performance waveguide components. In operation, the array will be cooled to 20 K and placed in the focal plane of a $\\sim 0.6$ meter telescope providing an angular resolution of  $\\sim1.5$ degrees. The LSPE experiment aims at large scale, high sensitivity measurements of CMB polarization, with multi-frequency deep measurements to optimize component separation. The STRIP Q-band channel is crucial to accurately measure and remove the synchrotron polarized component, while the W-band channel, together with a bolometric channel at the same frequency, provides a crucial cross-check for systematic effects. ",
        "introduction": "\\label{sec:intro}  %  \\input{introduction} ",
        "conclusions": "\\input{conclusions}"
    },
    "1208/1208.2161_arXiv.txt": {
        "abstract": "Adaptive optics observations of the flattened nuclear star cluster in the nearby edge-on spiral galaxy NGC~4244 using the Gemini Near-Infrared Integral Field Spectrograph (NIFS) have revealed clear rotation. Using these kinematics plus 2MASS photometry we construct a series of axisymmetric two-component particle dynamical models with our improved version of {\\sc nmagic}, a flexible $\\chi^2$-made-to-measure code. The models consist of a nuclear cluster disc embedded within a spheroidal particle population.  We find a mass for the nuclear star cluster of $\\mathrm{M}=1.6^{+0.5}_{-0.2}\\times 10^7 \\Msun$ within $\\sim 42.4$~pc ($2\\arcsec$).  We also explore the presence of an intermediate mass black hole and show that models with a black hole as massive as $\\Mbh = 5.0 \\times 10^5~\\Msun$ are consistent with the available data.  Regardless of whether a black hole is present or not, the nuclear cluster is vertically anisotropic ($\\beta_z < 0$), as was found with earlier two-integral models. We then use the models as initial conditions for $N$-body simulations.  These simulations show that the nuclear star cluster is stable against non-axisymmetric perturbations.  We also explore the effect of the nuclear cluster accreting star clusters at various inclinations.  Accretion of a star cluster with mass $13\\%$ that of the nuclear cluster is already enough to destroy the vertical anisotropy, regardless of orbital inclination. ",
        "introduction": "\\label{sec:intro}  Studies of the centres of galaxies across the Hubble sequence have shown that they frequently host central massive objects such as massive nuclear star clusters (NCs) and supermassive black holes (SMBHs).  NCs are present in roughly 75\\% of low and intermediate luminosity disc and elliptical galaxies \\citep{boe_etal_02, cot_etal_acsvcs8_06}.  These NCs are intrinsically very luminous, with typical $M_I \\sim -12$, and sizes similar to globular clusters \\citep[$r_{eff} \\sim 5$pc;][]{boker04a}.  Two hypotheses have been offered to explain NC formation.  One scenario envisages NCs forming in situ out of gas cooling onto the centre \\citep{milosa_04, bek_etal_06, bekki_07}.  Alternatively, NCs may form from star clusters merging at the centres of galaxies \\citep{tre_etal_75, lot_etal_01, cap-dol_mio_08, agarwal+2011, antonini+2012a, antonini2012b}. Which hypothesis is correct determines whether NC growth is limited by the supply of star clusters from the host galaxy \\citep{antonini2012b} or regulated by feedback from in situ star formation \\citep{mcl_etal_06}.  The assembly history of NCs can be constrained from their morphology, stellar populations and kinematics.  In late-type spirals, NCs have been found to consist of multiple stellar populations, typically a young population ($<100$ Myr), and a dominant population older than 1 Gyr \\citep{dav_cou_02, sch_etal_03, rossa+06, wal_etal_06}.  The {\\it Hubble Space Telescope} has revealed that the NCs of edge-on galaxies host multiple stellar populations associated with different morphological components \\citep{seth06}.  These NCs consist of young blue nuclear cluster discs (NCD) and older nuclear cluster spheroids (NCS).  Optical spectra of the NC in the edge-on Scd galaxy NGC~4244 ($i\\approx90\\degrees$), the nearest galaxy in the sample of \\citet{seth06} \\citep[D=4.37~Mpc;][]{seth05a}, indicates the presence of multiple stellar populations, while near infrared spectroscopy showed that the NC is rapidly rotating \\citep{seth+08b}. Using $N$-body simulations \\citet[][hereafter H11]{hart_etal_11} showed that the NC in NGC 4244 cannot have assembled more than half its mass via the accretion of star clusters.  NCs exhibit several scaling relations.  The luminosity of NCs correlates with that of their host galaxy \\citep{boe_etal_02, cot_etal_acsvcs8_06, erwin+gadotti10}.  A number of studies also found that their mass, \\Mnc, correlates with the velocity dispersion of the host bulge, the \\Mncsig\\ relation \\citep{fer_etal_06, weh_har_06, rossa+06}.  Early work found that this \\Mncsig\\ relation is parallel to the \\Msig\\ relation of SMBHs \\citep{geb_etal_00, fer_mer_00}, with NCs being about $10\\times$ more massive, at the same $\\sigma$, as SMBHs.  However recent work has questioned how comparable NCs and SMBHs are.  \\citet{erwin_etal_2012} find that NCs and SMBHs follow different relations, with SMBH masses correlated with properties of the bulge, while NCs seem to correlate better with properties of the entire host galaxy. Instead, both \\citet{leigh+12} and \\citet{scott_graham12} show that there is an $\\Mnc-\\sigma$ relation but with a significantly different slope than for SMBHs. It is not clear at present whether these differences are intrinsic to NC and SMBH growth or whether they are due to the fact that the scaling relations depend on Hubble type.  In particular, some recent studies have suggested that SMBHs and NCs  in late-type galaxies do not follow the same scaling relations as in early-types \\citep{greene+2010, erwin+gadotti12}. Some galaxies host both a NC and a SMBH \\citep{seth08, graham+spitler_09b}. The relative properties of NCs and SMBHs in such galaxies could constrain the relationship between these objects.  For instance, by constructing an $(\\Mbh+\\Mnc)-\\sigma$ relation which includes the mass of both the NC and of the SMBH \\citep{graham_etal_11}, \\citet{graham12} found a flatter relation than the \\Msig\\ relation.  But the small existing sample of objects with known NCs and SMBHs is currently too small to obtain a clear picture \\citep[e.g.][]{neumayer_etal_12}.  Progress in determining whether NCs and SMBHs are related therefore requires improving the statistics of such measurements.  Moreover, a better understanding of the mass assembly of NCs in late-type galaxies is vital.  It is generally thought that AGN feedback is responsible for the \\Msig\\ relation \\citep[e.g.,][]{sil_ree_98, aking_03, mur_etal_05, dim_etal_05, saz_etal_05, spr_etal_05b, joh_etal_09}, but scenarios where this relation arises because the galaxy regulates SMBH growth \\citep[e.g.][]{bur_sil_01, kaz_etal_05, mir_kol_05} or purely indirectly by the hierarchical assembly through galaxy merging \\citep{hae_kau_00, ada_etal_01, ada_etal_03, jahnke_maccio10} have also been proposed. If gas inflow plays a more important role in the growth of NCs then this opens the possibility that some form of feedback drives the scaling relations in both SMBHs and NCs \\citep[e.g.][]{mcl_etal_06}.  To help shed light on the formation of NCs in late-type galaxies, in this paper we study the NC in the nearby Sc galaxy NGC 4244. H11 modelled this NC using two-integral {\\sc jam} models \\citep{cappellari08}, obtaining a mass of $(1.1\\pm 0.2) \\times 10^7~\\Msun$.  In this paper we build three-integral particle models of the same NC and use them as initial conditions for $N$-body simulations to explore its sensitivity to star cluster accretion. The outline of this paper is as follows.  Section~\\ref{sec:obsdata} describes the observational data and how they are used in the dynamical modelling.  Our modelling method, the $\\chi^2$-M2M code {\\sc nmagic} based on \\citet{delo+07, delo+08, delo+09}, is described in Section~\\ref{sec:methods} including additional code development.  We construct various axisymmetric particle models of the nuclear region of NGC~4244 in Section~\\ref{sec:models} using this improved code.  The models consist of a NCD and NCS having separate mass-to-light ($M/L$) ratios.  Using the best model as initial conditions for $N$-body simulations, we explore the evolution of the NC in Section~\\ref{sec:nbody}.  Section~\\ref{sec:conclusions} discusses our results in the context of NC formation. ",
        "conclusions": "\\label{sec:conclusions}  We have performed a dynamical study of the nuclear cluster in the edge-on spiral galaxy NGC~4244 taking into account different morphological components, which are the galaxy main disc (MD), nuclear cluster disc (NCD) and nuclear cluster spheroid (NCS).  We have constructed axisymmetric dynamical particle models accounting for the MD, the NCD and the NCS. We find a total NCS mass $M_{\\rm NCS} = 1.6^{+0.5}_{-0.2}\\times 10^7~ \\Msun$ within approximately $42.4$~pc ($2\\arcsec$).  Both the fits of \\citet{seth05a} and of \\citet{fry+1999} show that there is no obvious bulge component in NGC 4244.  Using the {\\it 2MASS} Large Galaxy Atlas \\citep{jarrett+2003} $K$-band magnitude, the total luminosity of the galaxy is $3.2 \\times 10^9~ \\Lsun$, and thus the galaxy stellar mass is $~ 2 \\times 10^9 ~ \\Msun$.  Figure \\ref{fig:ferrarese} plots the NC mass compared with the $M_{CMO}-M_{gal}$ relation.  The NC sits above this relation.  \\begin{figure} \\centering \\vskip0.4truecm \\includegraphics[width=0.8\\hsize,angle=-90.0]{Mcmo-Mgal-cor.ps} \\caption[]{Mass of the NC of NGC 4244 versus the mass of the host galaxy compared with the $M_{CMO}-M_{gal}$ relations of \\citet{fer_etal_06}. The solid red and black lines show the correlations for NCs and SMBHs in early-type galaxies, respectively, with 1 $\\sigma$ confidence levels shown as dashed lines.  The blue lines show the relation of \\citet{scott_graham12}.} \\label{fig:ferrarese} \\end{figure}   \\subsection{Vertical anisotropy}  The kinematics are moderately tangentially anisotropic inside \\re\\ with an anisotropy parameter $\\beta_z \\sim -0.1$.  This is in good agreement with the {\\sc jam} models of H11.  H11 showed that $\\beta_z < 0$ requires high-inclination infall of star clusters onto a pre-existing nuclear cluster.  In our accretion simulations onto a more realistic model of the NC we found that even the accretion of a star cluster of just $13\\%$ the mass is enough to erase the vertical anisotropy.  This raises questions about whether such anisotropy can be due to accretion at all.  It also hints that, unless we are observing the NC of NGC 4244 at a special time, it cannot sustain accretion of $\\ga 10\\%$ mass as suggested by H11.  We therefore tested whether the assumption of a perfectly edge-on nuclear cluster may bias the modelled vertical anisotropy to negative values if the real inclination is somewhat smaller.  The smallest inclination at which we were able to deproject the NCD photometry was $83\\degrees$.  Using a model deprojected at this assumed inclination and the observed kinematics, we built {\\sc nmagic} models assuming $M/L_{\\mathrm NCS}=1.5$ starting from the best-fit edge-on model M2. The dashed black lines in Figure \\ref{fig:anisotropywithbh} show that the 2-D anisotropy, $\\beta_z = 1 - \\sigma_z^2/\\sigma_R^2$, and 3-D anisotropy, $B_z = 1 - 2\\sigma_z^2/(\\sigma_R^2 + \\sigma_\\phi^2)$, are barely changed compared to the edge-on case (solid black lines) and remain negative.  Thus a negative vertical anisotropy is not an artifact of assuming the nuclear cluster is perfectly edge-on.  We finally explore whether the recovered $\\beta_z$ changes if we include IMBHs in the models.  In Figure \\ref{fig:anisotropywithbh} we plot both the vertical anisotropies for varying $\\Mbh$.  While increasing $\\Mbh$ raises the vertical anisotropy, it still remains negative within \\re.  We conclude that the NC must be vertically anisotropic even if a black hole were present.   \\begin{figure} \\centering \\includegraphics[width=0.7\\hsize,angle=-90.0]{vertaniso2D.ps} \\includegraphics[width=0.7\\hsize,angle=-90.0]{vertaniso3D.ps} \\vskip0.2truecm \\caption[]{Profiles of 2-D (top) and 3-D (bottom) vertical anisotropy. The black, blue, green and red lines correspond to $\\Mbh = 0, 0.8, 1.23$, and $3.0 \\times 10^5~ \\Msun$, respectively.  The dashed lines correspond to assuming that the nuclear cluster is inclined at $83\\degrees$ instead of being perfectly edge-on.} \\label{fig:anisotropywithbh} \\end{figure}   \\subsection{Summary}  We have built dynamical models of the NC in the nearby, edge-on late-type galaxy NGC 4244. Using particle re-sampling, we were able to obtain a narrow distribution of weights in our {\\sc nmagic} models allowing us to use the models as initial conditions in $N$-body simulations. Our results can be summarised as follows:  \\begin{itemize}  \\item We find a mass of the spheroidal component of the NC, $M_\\mathrm{NCS} = 1.6^{+0.5}_{-0.2} \\times 10^7 ~ \\Msun$ within 42.4 pc.  The mass within 15 pc is $\\sim 1.0 \\times 10^7 \\Msun$, in very good agreement with the value estimated by \\citet{hart_etal_11} using two-integral {\\sc jam} models. This mass puts the nuclear cluster above the \\Mnc-$\\mathrm{M_{Gal}}$ relation.  \\item The mass of the bluer disc component of the nuclear cluster is less well constrained and covers the range $3.6 \\times 10^5 \\Msun \\la M_\\mathrm{NCD} \\la 14.4 \\times 10^5 \\Msun$.  \\item Our three-integral models are consistent with no black hole as well as with a black hole as massive as $4.6 \\times 10^5 ~ \\Msun$.  This upper limit is larger than the one allowed by two-integral {\\sc jam} models.  \\item Simulations show that the model without a black hole is stable against axisymmetric perturbations.  This stability derives from the large Toomre-$Q$ of the system.  \\item Regardless of whether a black hole is present or not, and of whether the nuclear cluster is perfectly edge-on or not, $\\beta_z$ and $B_z$ are both negative.  Accretion of a star cluster of as little as $13\\%$ by mass is enough to drive $\\beta_z$ to positive values, regardless of the orbital geometry.  It remains unclear, therefore, how $\\beta_z < 0$ arose.  \\end{itemize}    \\bigskip \\noindent"
    },
    "1208/1208.2482_arXiv.txt": {
        "abstract": "In `entropic cosmology', instead of a cosmological constant $\\Lambda$, an extra driving term is added to the Friedmann equation and the acceleration equation, taking into account the entropy and the temperature on the horizon of the universe. By means of the modified Friedmann and acceleration equations, we examine a non-adiabatic-like accelerated expansion of the universe in entropic cosmology. In this study, we consider a homogeneous, isotropic, and spatially flat universe, focusing on the single-fluid (single-component) dominated universe at late-times. To examine the properties of the late universe, we solve the modified Friedmann and acceleration equations, neglecting high-order corrections for the early universe. We derive the continuity (conservation) equation from the first law of thermodynamics, assuming non-adiabatic expansion caused by the entropy and temperature on the horizon. Using the continuity equation, we formulate the generalized Friedmann and acceleration equations, and propose a simple model. Through the luminosity distance, it is demonstrated that the simple model agrees well with both the observed accelerated expansion of the universe and a fine-tuned standard $\\Lambda$CDM (lambda cold dark matter) model. However, we find that the increase of the entropy for the simple model is likely uniform, while the increase of the entropy for the standard $\\Lambda$CDM model tends to be gradually slow especially after the present time. In other words, the simple model predicts that the present time is not a special time, unlike for the prediction of the standard $\\Lambda$CDM model. ",
        "introduction": "Numerous cosmological observations have implied a new paradigm for the cosmic expansion history, i.e., an accelerated expansion of the universe \\cite{Tyson1988,PERL1998a,PERL1998b,Riess1998,Riess2004,Riess2007,SN1,Tegmark1,Spergel1,Wood2007,Kowalski1,Hicken1,EKomatsu1,WMAP2011}. To explain the accelerated expansion, an additional energy component called `dark energy' is usually added to both the Friedmann equation and the Friedmann--Lema\\^{i}tre acceleration equation, where general relativity is assumed to be correct. In particular, $\\Lambda$CDM models, which assume cold dark matter (CDM) and a cosmological constant $\\Lambda$, have been suggested as an elegant description of accelerated expansion \\cite{Fukugita01,Carroll01,Ryden1,Hartle1,Sato1,Weinberg1,Roy1}. (For other models, see Refs.\\ \\cite{Weinberg1,Roy1,Miao1,Bamba1} and references therein.)  Recently, Easson, Frampton, and Smoot \\cite{Easson1,Easson2} have proposed that an extra driving term  should be added to the Friedmann--Lema\\^{i}tre acceleration equation. The additional entropic-force term can explain the accelerated expansion of the late universe \\cite{Easson1} and the inflation of the early universe \\cite{Easson2}, without introducing new fields \\cite{Koivisto1}. In the entopic-force scenario, called `entropic cosmology', the additional driving term is derived from the usually neglected surface terms on the horizon of the universe in the gravitational action, assuming that the horizon has an entropy and a temperature \\cite{Easson1}. (In fact, the entropy and temperature are related to the Bekenstein entropy \\cite{Bekenstein1} and the Hawking temperature \\cite{Hawking1} of black holes on an event horizon.) Since then, many researchers have extensively examined entropic cosmology from various viewpoints \\cite{Koivisto1,Cai1,Cai2,Qiu1,Casadio1,Danielsson1,Myung1,Costa1,Basilakos1}. (The possibility that the entropic force on the horizon can explain the accelerating universe \\cite{Easson1} should be distinguished from the idea that gravity itself is an entropic force \\cite{Padma1,Verlinde1}.)  In entropic cosmology, since the entropy on the horizon is assumed, the entropy can increase during the evolution of the universe. Therefore, it is possible to consider that the evolution of the universe is a kind of non-adiabatic process, unlike in standard cosmology, in which an adiabatic (isentropic) expansion is assumed. Nevertheless, such a non-adiabatic-like expansion of the universe has not yet been extensively investigated in entropic cosmology and has been considered in only a few studies \\cite{Cai1,Cai2, Qiu1,Casadio1,Danielsson1}. Therefore, it is important to examine the non-adiabatic-like (hereafter non-adiabatic) process, to acquire a deeper understanding of entropic cosmology, especially from a thermodynamics viewpoint. Also, after the discovery of black hole thermodynamics \\cite{Bekenstein1,Hawking1}, the entropy of the universe was examined by many researchers \\cite{Davies1,Davies2,Davies3,Frautsch1,Barrow1,Barrow2,Sugimoto1}. In particular, since the late 1990s, the entropy of the universe has been extensively discussed in a universe undergoing accelerated expansion \\cite{Easther1,Barrow3,Davies11,Davis01,Davis00,Jaeckel01,Gong00,Gong01,Setare1,Cline1}. However, evolution of the entropy has not been studied in entropic cosmology, although entropy plays an important role.  In this context, we examine a non-adiabatic expansion of the universe and discuss the evolution of the entropy in entropic cosmology. For this purpose, we derive the continuity (conservation) equation from the first law of thermodynamics, taking into account the non-adiabatic process caused by the entropy and the temperature on the horizon. If the modified Friedmann and Friedmann--Lema\\^{i}tre acceleration equations are used, the continuity equation can be derived from the two equations without using the first law of thermodynamics. This is because two of the three equations are independent \\cite{Ryden1}. However, in this study, we derive the continuity equation from the first law of thermodynamics, since the first law is the fundamental conservation law. Using the obtained continuity equation, we formulate the generalized Friedmann and Friedmann--Lema\\^{i}tre acceleration equations. In addition, we propose a simple model based on the formulation. It should be noted that we do not discuss entropic inflation \\cite{Easson2,Cai2} in the early universe, since we focus on the late universe \\cite{Easson1,Gao1} to examine the fundamental properties of the universe in entropic cosmology.  The present paper is organized as follows. In Sec. II, we give a brief review of the two modified Friedmann equations in entropic cosmology. In this section, we examine the properties of the single-fluid dominated universe. In Sec. III, we derive the modified continuity equation from the first law of thermodynamics, assuming a non-adiabatic expansion of the universe. We also discuss generalized formulations of entropic cosmology and propose a simple model. In Sec. IV, we compare the simple model with the observed supernova data and several $\\Lambda$CDM models. Finally, in Sec. V, we present our conclusions. ",
        "conclusions": "We have examined non-adiabatic expansion of the late universe and discussed the evolution of the entropy on the Hubble horizon, to study entropic cosmology from a thermodynamics viewpoint. For this purpose, we have employed the two modified Friedmann equations, i.e., the modified Friedmann equation and the modified acceleration equation. First of all, based on the two equations for entropic cosmology, we have examined the properties of the single-fluid dominated universe, neglecting high-order terms for quantum corrections. Consequently, we can systematically summarize the properties of the late universe, through a parameter $C_{1}$ related to entropic-force terms. It is found that, at late times, the entropy on the Hubble horizon increases slowly with decreasing $C_{1}$ (or as the influence of entropic-force terms increases), while the expansion of the universe accelerates.  We have also derived the continuity equation from the first law of thermodynamics, assuming non-adiabatic expansion of the universe. Using the obtained continuity equation, we have formulated the generalized Friedmann and acceleration equations, and have proposed a simple model as a possible model. Through the luminosity distance, it is successfully shown that the simple model can explain the present accelerating universe and agrees well with both the supernova data and the fine-tuned standard $\\Lambda$CDM model. On the other hand, the increase of the entropy for the simple model is uniform, although the increase of the entropy for the standard $\\Lambda$CDM model is gradually slow especially after the present time. In other words, the simple model implies that the present time is not a special time, unlike for the prediction of the standard $\\Lambda$CDM model. We find that the present simple model predicts another future which is different from the standard $\\Lambda$CDM model.  The present study has revealed the fundamental properties of the non-adiabatic expanding universe in entropic cosmology. As one of several possible scenarios, the generalized formulation and the simple model considered here will help in understanding the accelerating expanding universe. Of course, it is difficult to determine $\\gamma$ related to the temperature on the Hubble horizon. However, at least in principle, it is possible to discuss the accelerating universe quantitatively, by means of the present entropic cosmology. Further discussions and observation data will be required to examine the present and future of the universe.  The modified continuity equation examined here has the so-called non-zero term on the right-hand side. Therefore, through the present paper, we call this the non-adiabatic process. If we employ an effective description similar to bulk viscous cosmology, the non-zero term on the right-hand side can be cancelled in appearance. Alternatively, the non-zero term may be interpreted as the interchange of energy between the bulk (the universe) and the boundary (the horizon of the universe). Accordingly, it will be necessary to study entropic cosmology from various viewpoints.       \\appendix"
    },
    "1208/1208.3092_arXiv.txt": {
        "abstract": "{} {We have obtained high contrast images of four nearby, faint, and very low mass objects 2MASS\\,J04351455-1414468, SDSS\\,J044337.61+000205.1, 2MASS\\,J06085283-2753583 and 2MASS\\,J06524851-5741376 (here after 2MASS0435-14, SDSS0443+00, 2MASS0608-27 and 2MASS0652-57), identified in the field as probable isolated young brown dwarfs.  Our goal was to search for binary companions down to the planetary mass regime.} {We used the NAOS-CONICA adaptive optics instrument (NACO) and its unique capability to sense the wavefront in the near-infrared to acquire sharp images of the four systems in $K_s$, with a field of view of $28~\\!''\\times28~\\!''$. Additional $J$ and $L'$ imaging and follow-up observations at a second epoch were obtained for 2MASS0652-57.} {With a typical contrast $\\Delta K_s=4.0-7.0$~mag, our observations are sensitive down to the planetary mass regime considering a minimum age of 10 to 120~Myr for these systems. No additional point sources are detected in the environment of 2MASS0435-14, SDSS0443+00 and 2MASS0608-27 between $0.1-12~\\!''$ (i.e about 2 to 250~AU at 20~pc). 2MASS0652-57 is resolved as a $\\sim230$~mas binary. Follow-up observations reject a background contaminate, resolve the orbital motion of the pair, and confirm with high confidence that the system is physically bound. The $J$, $K_s$ and $L'$ photometry suggest a $q\\sim0.7-0.8$ mass ratio binary with a probable semi-major axis of 5-6~AU. Among the four systems, 2MASS0652-57 is probably the less constrained in terms of age determination. Further analysis would be necessary to confirm its youth. It would then be interesting to determine its orbital and physical properties to derive the system's dynamical mass and to test evolutionary model predictions.} {} ",
        "introduction": "The statistical properties of low-mass star and brown dwarf multiples set stringent constraints on star-formation theories (see Duchene et al. 2007 for a review).  Multiplicity frequency, mass ratio and separation distributions can be compared between star forming regions (SFR's) of various ages and densities and the older field population. Direct imaging surveys of very low mass objects in the field yield a binary frequency of 20-30\\% for M dwarfs (Marchal et al. 2003; Janson et al. 2012) and 15\\% for L and T dwarfs (e.g., Bouy et al. 2003; Burgasser et al. 2003). Among young systems, Ahmic et al. (2007) and Biller et al. (2011) derive a binary fraction of less than 11\\% and 9\\% in the Chamaeleon I and the Upper Sco regions respectively. They both confirm the trend observed in the field of a mass-dependency of the binary frequency. A higher multiplicity rate in SFR's compared with the field, as for T Tauri stars, is not seen for young late-type M dwarfs.  However the multiplicity properties are likely to be different as evidenced by the discovery of a population of wide ($>15$~AU) brown dwarf binaries in young, nearby clusters (Chauvin et al. 2004; Jayawardhana \\& Ivanov 2006; B\\'ejar et al. 2008; Todorov et al. 2010). Individual, young, and tight binaries are particularly important for a direct determination of the dynamical mass to calibrate theoretical masses derived from evolutionary models (Mathieu et al. 2007; Bonnefoy et al. 2009). \\begin{table*}[t] \\caption{Description of the target properties} \\begin{center} \\small \\begin{tabular}{lllllllllll}     % \\noalign{\\smallskip}\\hline \\noalign{\\smallskip}\\hline  \\noalign{\\smallskip} Name            & $\\alpha$        &  $\\delta$       &    $\\mu_{\\alpha}^a$   & $\\mu_{\\delta}^a$     & Vrad       & SpT      & $d^a$     &   J     & K        & Ref.$^b$ \\\\ & [J2000]         & [J2000]         &    (mas/yr)         & (mas/yr)       & (km/s)     &          & (pc)  &   (mag) & (mag)    &            \\\\ \\noalign{\\smallskip}\\hline                  \\noalign{\\smallskip} 2MASS0435-14\t& 04 35 14.6      & -14 14 47       & $9\\pm14$            & $16\\pm14$      &            &   M6$\\delta\\pm$1     & $8.6\\pm1.0$       & 11.88  & 9.95    & 1 \\\\ SDSS0443+00\t& 04 43 37.6      & +00 02 05       & $28\\pm14$ \t  & $-99\\pm14$     &            & M9$\\gamma$      & $16.2\\pm2.1$      & 12.51  & 11.22   & 1, 2, 3, 4\\\\ 2MASS0608-27\t& 06 08 52.8      & -27 53 58       & $8.9\\pm3.5$         & $10.7\\pm3.5$   & $24\\pm1$   & M8.5$\\gamma$     & $31.2_{-3.2}^{+4.0}$ & 13.59  & 12.37   & 1, 5, 6 \\\\ 2MASS0652-57\t& 06 52 48.5      & -57 41 38       & $0.1\\pm3.4$ \t  & $29.2\\pm3.3$   &            & M8$\\beta$       & $31.9_{-2.9}^{+3.7}$ & 13.63  & 12.45   & 6, 7\\\\ \\noalign{\\smallskip}\\hline  \\noalign{\\smallskip} \\end{tabular} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Distances derived from spectrophotometry with M$_J$ estimated from the spectral type/M$_J$ calibration from Cruz et al. (2003) for 2MASS0435-14 and SDSS0443+00. Proper motion and parallax measurements from Faherty et al. (2012) for 2MASS0608-27 and 2MASS0652-57 \\item[$^{\\mathrm{b}}$] References: (1) Cruz et al. 2003, AJ, 126, 2421, (2) Cruz et al. 2007, AJ, 133, 43, (3) Reid et al. 2008, 136, 1290, (4) Reiner \\& Basri 2009, AJ, 705, 1416, (5) Rice et al. 2010, ApJ, 715, 165, (6) Faherty et al. (2012) and (7) Reid et al. (2008) \\end{list} \\end{center} \\end{table*}  For very low mass binaries with high mass ratio, the mass of the binary companion can enter the planetary mass regime. The origin of such a population of planetary mass companions (PMCs) can be hard to infer as the stellar and planetary formation mechanisms probably overlap. From radial velocity surveys, the observed frequency of giant planets around M dwarfs at small separation ($<3$~AU) is relatively small ($f < 1-2~\\%$; Bonfils et al. 2011) compared with solar-type stars ($f < 6-9~\\%$; Udry \\& Santos 2007) and could indicate a mass dependency of the core accretion mechanism efficiency to form giant planets. At larger separations, the situation is less clear. We may expect that alternative mechanisms to core accretion such as cloud or disk fragmentation may form a population of planetary mass companions such as 2M1207\\,b (Chauvin et al. 2004). Consequently, despite the fact that most exoplanet imaging surveys are now biased towards young, intermediate mass stars, the search for PMCs around low-mass stars and brown dwarfs remains important to understand how planetary formation evolves with the stellar mass and the distance to the star (Delorme et al. 2012).  In the course of our deep imaging survey of 88 young, nearby stars with NACO at VLT (Chauvin et al. 2010), an additional sub-sample of four probable intermediate-young brown dwarfs were observed taking advantage of the infrared (IR) wavefront sensing system of the NACO adaptive optics (AO) instrument. One is a recently confirmed member of the $\\beta$ Pictoris moving group (Rice et al. 2010) while the other 3 are low-gravity M dwarfs indicating a likely age lower than a few hundred Myr. We report, in section 2, a summary of the target properties. In section 3, we describe our observations, including the instrument setup and the atmospheric conditions. In section 4, we present the results of this imaging campaign, in terms of detection limits. We also report that 2MASS0652-57 is a $\\sim230$~mas binary. Observations at 2 epochs enable us to resolve the system orbital motion. Finally, we briefly discuss the status of this young binary, and of its membership to any known young moving groups. ",
        "conclusions": ""
    },
    "1208/1208.1097_arXiv.txt": {
        "abstract": "We report one of several homologous non-radial eruptions from NOAA active region (AR) 11158 that are strongly modulated by the local magnetic field as observed with the \\textit{Solar Dynamic Observatory} (\\textit{SDO}). A small bipole emerged in the sunspot complex and subsequently created a quadrupolar flux system. Non-linear force-free field (NLFFF) extrapolation from vector magnetograms reveals its energetic nature: the fast-shearing bipole accumulated $\\sim$2$\\times$10$^{31}$ erg free energy (10$\\%$ of AR total) over just one day despite its relatively small magnetic flux (5$\\%$ of AR total). During the eruption, the ejected plasma followed a highly inclined trajectory, over 60$^\\circ$ with respect to the radial direction, forming a jet-like, inverted-Y shaped structure in its wake. Field extrapolation suggests complicated magnetic connectivity with a coronal null point, which is favorable of reconnection between different flux components in the quadrupolar system. Indeed, multiple pairs of flare ribbons brightened simultaneously, and coronal reconnection signatures appeared near the inferred null. Part of the magnetic setting resembles that of a blowout-type jet; the observed inverted-Y structure likely outlines the open field lines along the separatrix surface. Owing to the asymmetrical photospheric flux distribution, the confining magnetic pressure decreases much faster horizontally than upward. This special field geometry likely guided the non-radial eruption during its initial stage. ",
        "introduction": "\\label{sec:intro}  Solar eruptive events derive their energy from the non-potential coronal magnetic field \\citep{forbes2000,hudson2011}. Reconnection takes place locally where the field gradient is large, but can alter the larger-scale field topology rapidly. The dissipated energy from the relaxing field accelerates particles, produces radiation, and heats and ejects plasma into the interplanetary space as a coronal mass ejection (CME).  Prior to eruption, energy builds up in the corona through flux emergence and displacement, which may take up to a couple of days \\citep{schrijver2009}. The slow evolution can be approximated by a series of quasi-stationary, force-free states in the low plasma-$\\beta$ coronal environment. This allows the estimation of AR energetics in non-flaring states, thanks to recent advances in photospheric field measurement and field extrapolation algorithms \\citep{regnier2006,thalmann2008,jing2009,sun2012}.  Besides the gross energy budget, the detailed magnetic configuration also proves important to the initiation, geometry, and scale of eruptions. In the case of a coronal jet, the direction of the ambient field (horizontal or oblique) directly determines the direction of the jet and its distinct emission features \\citep[two-sided or ``anemone'' type,][]{shibata1997}. Observation and modeling demonstrate that the overlying field provides a critical constraint on CME's speed and trajectory \\citep{liuyang2007,gopalswamy2009,wangyuming2011}.  Theoretical studies have extensively explored the role of topological features in reconnection \\citep{demoulin1996,priest2000,longcope2005}. Their applications to solar events usually involved the results of potential or linear force-free field extrapolation \\citep{aulanier2000,fletcher2001,mandrini2006}, or magnetohydrodynamic (MHD) simulations that qualitatively reproduce the observed phenomena \\citep{moreno2008,pariat2009,masson2009,torok2011}.  Here we report one of several similar non-radial eruptions that are strongly modulated by the local magnetic field as observed with the \\textit{Solar Dynamic Observatory} (\\textit{SDO}). Using vector magnetograms from the Helioseismic and Magnetic Imager (HMI) \\citep{schou2012,hoeksema2012} aboard \\textit{SDO} and a non-linear force-free field (NLFFF) extrapolation, we monitor the AR evolution and explain the magnetic topology that leads to the curious features during the eruption. The Atmospheric Imaging Assembly \\citep[AIA;][]{lemen2012} and other observatories recorded these features and provide guidance for our interpretation.  In Section~\\ref{sec:method} we briefly describe the data and the extrapolation algorithm. We first present observations of the eruption in Section~\\ref{sec:erupt}, and then come back in Section~\\ref{sec:evo} to explain the magnetic field and energy evolution leading to the event. In Section~\\ref{sec:topo}, we interpret this curious event based on the magnetic field topology. We discuss the interpretation in Section~\\ref{sec:discuss} and summarize in Section~\\ref{sec:summary}.   \\begin{figure}[t!] \\centerline{\\includegraphics[width=3in]{homolog.eps}} \\caption{Full-disk, unsharp masked AIA 171 {\\AA} image at 17:28:15 UT on February 14, 2011 showing the non-radial eruption. Inset shows the enhanced image of the ejecta. The two flux-rope-like structures with a shared eastern footpoint are marked as FR1 and FR2. Animation of a 20-hr interval shows at least five similar eruptions. (An animation of this figure is available at \\url{http://sun.stanford.edu/~xudong/Article/Cusp/homolog.mp4}.) \\label{f:homolog}} \\end{figure}  \\begin{figure*}[t!] \\centerline{\\includegraphics{ejecta.eps}} \\caption{Geometry of the non-radial eruption. O and $\\rm{{O'}}$ mark the eruption site. (a) SECCHI EUVI 195 {\\AA} image from \\textit{STEREO}-A, about 87$^\\circ$ ahead of \\textit{SDO}. Due to the tilt of the solar rotational axis, the \\textit{SDO} and \\textit{STEREO} north are offset by 6.8$^\\circ$. (b) AIA 193 {\\AA} image of the same ejecta, taken 4 s later than (a). The projected N-S length of the ejecta ($|{\\rm{OR}}|$) is identical to that in (a) ($|\\rm{{O'R'}}|$), where ${\\rm{OR}}$ and ${\\rm{O'R'}}$ represent the projection of line segments OP and OQ in the N-S direction in \\textit{SDO}'s plane-of-sky, respectively. The scales of (a) and (b) are different in order to better show the features of interest. (c) Schematic diagram explaining the determination of the ejecta's geometry. SDO's west, north, and LOS directions are taken as $x$, $y$, and $z$ axis. The pink arrow represents the ejecta, its projected shape viewed from EUVI and AIA are shown as pink dashed lines on green and brown planes. The local radial vector is about W13S04 to LOS. The inclination $\\delta$ is about 43$^\\circ$; azimuth $\\alpha$ about 34$^\\circ$. See Section~\\ref{sec:erupt} for details. (d) AIA 171 {\\AA} image of the post-eruption AR; Y marks the top of the cusp and the base of the jet. The boxed region is used to construct panel (e). Purple/pink contours are for HMI LOS field at $\\pm$200 G. (e) Space-time diagram showing the speed of ejecta and jet. Three dashed lines (starting near 17:25, 17:36, and 17:49 UT) indicate a projected speed of 500, 330 and 280 km s$^{-1}$, respectively. Panels (a), (b), and (d) are displayed in a square-root scale. (An animation of this figure is available at \\url{http://sun.stanford.edu/~xudong/Article/Cusp/ejecta.mp4}.) \\label{f:ejecta}} \\end{figure*} ",
        "conclusions": "\\label{sec:discuss}  \\subsection{On the Coronal Field Topology} \\label{subsec:d_topo}  How common is the magnetic topology determined here? A previous study focused on the quadrupolar configuration of AR 10486 during the 2003 X-17 flare \\citep{mandrini2006}. The major eruption was found to involve reconnection at the quasi-separatrix layers \\citep[QSL;][]{demoulin1996}, while a smaller brightening was associated with a similar coronal null point determined using a linear force-free extrapolation. In another quadrupolar region AR 11183, similar cusp and jet structures existed at a much larger scale \\citep{filippov2012}. The white-light jet extended over multiple solar radii.  We analyze the entire 36-hr series, searching for consistency in time. The coronal null at 9 Mm appeared in a few frames early on February 14, distinct from all other candidates which were mostly below 4 Mm in weak field regions. Starting from 15:35 UT, it appeared at a nearly constant location (within 3 Mm of the first detected null) in over half the frames afterwards (22/42, until February 15 00:00 UT), while the near-surface nulls rarely repeated in two consecutive time steps. We have applied a different null-searching method based on the Poincar\\'{e} index theorem \\citep{greene1992} and found similar results (23/42, 20 identical to the trilinear method, with 3 additional and 2 missed detections). The repeated detection of null points and the observed homologous eruptions (Figure~\\ref{f:energy}(d)) suggest the aforementioned topology is characteristic for this quadrupolar system.  We compute at 1-hr cadence the ``squashing factor'' $Q$ that describes the field mapping gradient \\citep{titov2002} by tracing individual field lines and measuring the differences between the two footpoint locations. High-$Q$ isosurface corresponds to QSLs. By inspecting the contour of $Q$ at different heights, we find that multiple QSL's tend to converge and intersect at about 9 Mm. Near the intersection, the field strength is weak, and the field line mapping gradient is invariably large, with or without null point. This illustrates the robustness of our interpretation despite the uncertainties in the extrapolation algorithm \\citep[e.g.][]{derosa2009} and the field measurement. (The uncertainties nevertheless can indeed affect the detailed fan-spine configuration, as discussed in Appendix~\\ref{a:skeleton}.)  We note that our PF extrapolation, with radial field as boundary condition and the Green's function method, does not detect any nulls above 5 Mm. Instead, we find a low-lying null at about 4 Mm in 13 frames, southwest to the NLFFF solution. The field configuration is less realistic, presumably because the current-free assumption does not agree with observation.  \\subsection{On the Flare Emissions} \\label{subsec:d_obs}  Owing to the LOS projection, the altitude of an on-disk HXR source cannot be unambiguously determined. We think S2 is a coronal source mainly because it appeared near the apex of cusp-shaped loops (P2/N2) which is typical for reconnecting field lines \\citep[e.g.][]{tsuneta1996}. In addition, its strong HXR emission (peak at $\\sim$60$\\%$ of the maximum) does not correspond to any bright flare ribbon. The closest chromospheric emission enhancement is a small patch (${\\rm{R_{P0}}}$ in Figure~\\ref{f:corona}(c)) within a fragmented positive flux about 5$\\arcsec$ to the east and south, whose intensity is much weaker than the ${\\rm{R_{P1}}}$/${\\rm{R_{N1}}}$ ribbons. This argues against the footpoint source interpretation.  We notice a dimmer, half-ring-like ribbon (${\\rm{R_{N0}}}$) farther north in the weak field area (Figure~\\ref{f:corona}(b)); both H$\\alpha$ (Figure~\\ref{f:corona}(c)) and EUV images (animation of Figure~\\ref{f:ejecta}) show its connection to P1. This structure is related to flux emerging into an encircling unipolar region \\citep[``anemone'' AR;][]{shibata1994}. Because the brightening ${\\rm{R_{N0}}}$ region possesses flux only a few percent of P1 \\citep[c.f.][]{reardon2011}, we consider this structure secondary. It does not affect our conclusions on the AR topology.  Because no HXR source was detected at the P2/N2 footpoints and the ${\\rm{R_{P2}}}$/${\\rm{R_{N2}}}$ ribbons were fainter than ${\\rm{R_{P1}}}$/${\\rm{R_{N1}}}$, we think the electrons primarily precipitated along the shorter P1/N1 loop during the flare. On the other hand, the P2/N2 loop produced much stronger SXR and EUV emission during the flare's late decaying phase. Almost 30 minutes later, SXR images (Figure~\\ref{f:corona}(f)) from the \\textit{Hinode} X-Ray Telescope \\citep[XRT;][]{golub2007} still showed a bright cusp structure above P2/N2.   \\begin{figure}[t!] \\centerline{\\includegraphics[width=2.6in]{jet.eps}} \\caption{Schematic illustration of the magnetic configuration and dynamics that may have led to the eruption. The structure resembles that of a blowout jet. The arcade (blue field lines above P1/N1) from the newly emerged bipole expands, reconnects with the pre-existing field (blue field lines from N2), becomes open (yellow field lines from N1), and the low-lying sheared/twisted core field (pink field lines between P1/N1) subsequently erupts. A possible initial reconnection site near is marked by the star; possible motions of the loops are denoted by thick arrows. Pre- and post- reconnection field lines are colored blue and yellow, respectively. The directions of the observed, \\textit{post-eruption} flow (Figure~\\ref{f:ejecta} and animation, see also \\cite{thompson2011,sujt2012}) are denoted by thin arrows. The inset shows the SXR difference image between 17:22:32 and 17:19:56 UT from \\textit{Hinode} XRT Ti Poly filter (FOV 72{\\arcsec}$\\times$60{\\arcsec}). The brightening P1/N2 loop is marked by a yellow circle; the brightening filament is visible in the foreground. \\label{f:jet}} \\end{figure}   \\subsection{On the Eruption Mechanism} \\label{subsec:d_mech}  When a bipole emerges, one leg of the new loop may reconnect with the oppositely directed, pre-existing open field. The released magnetic energy heats the plasma and produces field-collimated outward flows, known as the ``standard'' jet phenomenon \\citep{shibata1997}. When the emerging field is sheared or twisted, its core may subsequently erupt. Events in this sub-class have recently been described as ``blowout'' jets \\citep{moore2010}.  Can this event be explained by the jet models? We find the inferred magnetic structure here resembles the blowout type. Illustrated in Figure~\\ref{f:jet}, the newly emerged bipole (P1 and the north part of N1) hosts a twisted core field. We speculate that the increasing flux leads to the expansion of the arcade loops above, which reconnect with the open, negative-polarity field from N2. This process opens up the arcade loops and acts to promote the eventual eruption of the core field below. The jet model predicts the brightening of the reconnected P1/N2 loop, which is indeed observed in the SXR images (inset of Figure~\\ref{f:jet}). However, in contrast to the expected jet behavior, no outward flows are observed during this stage. The jet-like, inverted-Y structure appeared only \\textit{after} the core field eruption and the accompanying M-class flare.  Propagating brightness disturbances in the post-eruption inverted-Y structure have been interpreted as pulsed plasma flow \\citep{sujt2012,tian2012}. The upflow from the left leg diverges and flows in opposite directions, upward in the thin spire and downward in the right leg (Figure~\\ref{f:jet} and Animation of Figure~\\ref{f:ejecta}). The flow is most pronounced in cooler EUV wavebands (e.g. 171 {\\AA}, $\\sim$0.6 MK) and is absent in SXR images.  In the standard jet model, these collimated flows are produced and heated by reconnection. The relatively low temperature observed here suggests a low-altitude reconnection site with cooler plasma supply \\citep[c.f.][]{sujt2012}, rather than the one near the base of the spire higher in the corona. The detailed dynamics of this event require further investigation which is out of the scope of this work.    We summarize our findings as follows.  \\begin{itemize} \\renewcommand{\\labelitemi}{$-$}  \\item Bipole emergence and shearing in a pre-existing sunspot complex introduced a large amount of free energy, despite its small flux. The new flux powered a series of homologous, non-radial eruptions.  \\item One typical eruption had an inclined trajectory about 66$^\\circ$ with respect to the radial direction. An inverted-Y structure consisted of cusp and jet formed in the wake of the eruption.  \\item The bipole emergence created an asymmetrical quadrupolar flux system. Field extrapolation suggests that the consequent, inclined overlying loops and the anisotropic magnetic pressure are responsible for the non-radial eruption.  \\item Extrapolation suggests a coronal null point at about 9 Mm, slightly below the apexes of the cusp-like loops. Its location is favorable for reconnection between different flux components in the quadrupolar system. The observed inverted-Y structure is likely related to the open negative field lines in part outlining the separatrix surface.  \\item Multiple flare ribbons brightened simultaneously during the accompanying flare. A coronal HXR source appeared near the inferred null point. These observations support our interpretation.  \\item The inferred magnetic structure resembles that of a blowout-type jet. Some observed features fit in the jet model, while others remain difficult to explain.  \\end{itemize}  The event studied here demonstrates the importance of detailed magnetic field topology during solar eruptions. Flux emergence in suitable environment can lead to fundamental changes in the coronal field geometry, which then place strong constraints on the plasma dynamics."
    },
    "1208/1208.1268_arXiv.txt": {
        "abstract": "We report homogeneous spectroscopic determinations of the effective temperature, metallicity, and projected rotational velocity for the host stars of 56 transiting planets. Our analysis is based primarily on the Stellar Parameter Classification (SPC) technique. We investigate systematic errors by examining subsets of the data with two other methods that have often been used in previous studies (SME and MOOG). The SPC and SME results, both based on comparisons between synthetic spectra and actual spectra, show strong correlations between $T_{\\rm eff}$, [Fe/H], and $\\log g$ when solving for all three quantities simultaneously. In contrast the MOOG results, based on a more traditional curve-of-growth approach, show no such correlations. To combat the correlations and improve the accuracy of the temperatures and metallicities, we repeat the SPC analysis with a constraint on $\\log g$ based on the mean stellar density that can be derived from the analysis of the transit light curves.  Previous studies that have not taken advantage of this constraint have been subject to systematic errors in the stellar masses and radii of up to 20\\% and 10\\%, respectively, which can be larger than other observational uncertainties, and which also cause systematic errors in the planetary mass and radius. ",
        "introduction": "\\label{sec:introduction}  In recent years the number of extrasolar transiting planets has expanded considerably, with discoveries being made at a rapid pace both from the ground and increasingly also from space by the \\emph{CoRoT} and \\emph{Kepler} missions. With this large assembly of data, studies have begun to focus on examining patterns and correlations among the global properties of these planets and their parent stars, and what this can tell us about planet formation and evolution.  While the characteristics of some of these systems are very well known (e.g., HD\\,209458, HD\\,189733, TrES-1), those of others are much less well determined and have remained so since their discovery.  Our knowledge of the planetary properties depends critically on understanding the parent stars. This is because, for transiting systems, the light curves only give information on the size of the planet relative to that of the star ($R_p \\propto R_{\\star}$), and spectroscopic observations only reveal the mass of the planet if we know the mass of the star ($M_p \\propto M_{\\star}^{2/3}$). The stellar mass and radius, in turn, depend on other properties that can be gleaned from the stars' spectra such as the effective temperature ($T_{\\rm eff}$), surface gravity ($\\log g$), and chemical composition (commonly represented by [Fe/H]).  For many of the known transiting planet systems, follow-up photometric observations have been undertaken after the initial discovery of the planet, usually for the purpose of measuring the times of mid-transit and seeking departures from strict periodicity (transit timing variations) that may indicate the presence of additional bodies in the system. These new transit light curve observations have also served to improve the radius determinations in many cases. However, it is much less common for known transiting systems to be re-observed spectroscopically. As a result, our knowledge of the stellar properties is often limited by whatever information was reported in the discovery papers, which is sometimes preliminary. Inaccuracies in the stellar $T_{\\rm eff}$, $\\log g$, and [Fe/H] propagate through to the determination of the planetary properties, and may obscure correlations with other quantities and prevent us from gaining valuable insight into the nature of planets. To make matters worse, the methods of determining $T_{\\rm eff}$, $\\log g$, and [Fe/H] in the literature are highly inhomogeneous, as they have been carried out by many groups using different assumptions and methodologies.  In one of the few studies to redetermine spectroscopic properties for transiting planet hosts in a uniform way, \\cite{Ammler:09} obtained new temperatures, surface gravities, and metallicities for 13 systems based on new or existing spectra. They combined their determinations with those for 11 additional systems made by others using similar techniques, and compiled a list of 24 host stars with uniformly derived properties. Comparison with results for the same stars by other groups uncovered significant systematic differences in some cases, the causes of which are unknown.  One of the motivations for the present paper is to derive spectroscopic properties in a homogeneous manner for a much larger sample of more than 50 transiting planet hosts, and thereby reduce any dispersion in the stellar and planetary properties that is caused by the variety of methodologies used in the past. For this we make use of the Stellar Parameter Classification (SPC) technique introduced by \\cite{Buchhave:12}. We also wish to understand potential systematic errors in such determinations. In particular, it has been recognized for some time \\citep[e.g.,][]{Sozzetti:07, Holman:07, Winn:08} that surface gravities are often poorly constrained by spectral analyses. This is an unfortunate limitation, because the surface gravity would otherwise help to establish the luminosity of the star (and hence its radius) by placing it on the H-R diagram.  In cases for which the surface gravity is poorly constrained, the use of an external constraint on the luminosity becomes very important to allow accurate determinations of the mass and radius of the star. However, a detail that has usually been overlooked is that the determinations of other spectroscopic quantities such as the temperature and metallicity can \\emph{also} be affected by the poor constraint on gravity. This is because the uncertainties in $T_{\\rm eff}$ and [Fe/H] are strongly correlated with $\\log g$ in at least some of the commonly used analysis techniques.  This can be a significant source of systematic error. Therefore, a second goal of our work is to compare spectroscopic determinations for a subset of the sample obtained using two additional procedures widely employed in previous studies, and to investigate and quantify systematics errors stemming from the weakly constrained gravities.  Ultimately the stellar quantities that enter into the calculation of the planetary characteristics are the masses and radii inferred from $T_{\\rm eff}$, $\\log g$, and [Fe/H], as mentioned earlier. A third objective of the present paper is to study and quantify how errors in the spectroscopic properties propagate into the stellar masses and radii.  The paper is organized as follows. In Sect.~\\ref{sec:observations} we report our spectroscopic observations, which consist of new and archival spectra obtained with three different telescopes. Our spectroscopic analysis techniques are described in Sect.~\\ref{sec:analysis}. Our results with and without the application of external constraints on the surface gravity are presented in Sect.~\\ref{sec:results}. Also presented there are the final results from this work. These are then compared with work by others in Sect.~\\ref{sec:comparison}. In Sect.~\\ref{sec:impact} we study the impact of the different assumptions regarding $\\log g$ on the stellar masses and radii. We discuss our results in Sect.~\\ref{sec:discussion}, and end with a summary of our conclusions in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  Accurate knowledge of the properties of the host stars in transiting exoplanet systems is essential to derive accurate characteristics for the planets. Considerable efforts have been devoted to improving the light curves of newly discovered as well as previously known transiting systems, but relatively little attention has been paid to refining the spectroscopic properties of the stars.  Here we have derived new effective temperatures, metallicities, and projected rotational velocities for 56 transiting planet systems in a homogeneous manner, using the SPC technique.  These determinations bring a needed measure of uniformity to the growing collection of stellar and planetary properties that should facilitate the discovery of patterns and correlations that may provide valuable insight into the nature of planets.  A key aspect of our spectral analysis is the application of an external constraint on the surface gravity, based on accurate knowledge of the mean stellar density of the star, which comes directly from the light curve modeling. Because $\\log g$ is usually weakly constrained by the spectra, fixing it as we have done here prevents errors in $\\log g$ from biasing the temperatures and metallicities, and from affecting the inferred stellar masses and radii.  Those biases come mainly from strong correlations between $T_{\\rm eff}$, [Fe/H], and $\\log g$ that are present in unconstrained determinations ($\\log g$ free), not only when applying SPC, but also in the widely used SME procedure as implemented by \\cite{Valenti:05}. Both of these methods are based on spectral synthesis. We find that the correlations are much smaller with MOOG, which uses a more classical curve-of-growth approach.  We investigate the interagreement among the three spectroscopic techniques by applying SME and MOOG to subsets of our stars, and we show that the temperatures and metallicities are generally in good accord after application of the $\\log g$ constraint (the mean differences being well under 50\\,K and 0.1~dex, respectively). We do, however, detect some remaining systematic trends as a function of temperature and metallicity that are occasionally larger in some regimes.  Virtually all current studies of transiting planets make use of the mean stellar density as a luminosity indicator to derive the stellar properties, either through a comparison with model isochrones, or using empirical relations. We demonstrate that not using $\\rho_{\\star}$ can incur errors in mass of up to 20\\%, and errors in the radius as large as 100\\% in some cases. Even though such errors are now usually avoided, many authors still retain the temperatures and metallicities obtained from \\emph{unconstrained} spectroscopic analyses, i.e., without fixing $\\log g$ to the more accurate values based on the light curve modeling. We demonstrate that this practice can lead to residual biases in $M_{\\star}$ of up to 20\\%, and systematic errors in $R_{\\star}$ up to 10\\% for the hotter stars, which will propagate through to the planetary properties. Such errors can be larger than other observational uncertainties, and may explain part of radius anomaly of some of the inflated Jovian planets.  In order to avoid this, we advocate performing a second iteration on the spectroscopic analysis (particularly when using SPC or SME) that fixes $\\log g$ to the value inferred from the photometrically determined mean stellar density."
    },
    "1208/1208.1432_arXiv.txt": {
        "abstract": "Planets form in active protoplanetary disks that sustain stellar jets. Momentum loss from the jet system may excite the planets' orbital eccentricity and inclination (Namouni 2005, AJ 130, 280). Evaluating quantitatively the effects of such excitation requires a realistic modeling of the momentum loss profiles associated with stellar jets. In this work, we model linear momentum loss as a time-variable stochastic process that results in a zero mean stellar acceleration. Momentum loss may involve periodic or random polarity reversals. We characterize orbital excitation as a function of the variability timescale and  identify a novel excitation resonance between a planet's orbital period and the jet's variability timescale where the former equals twice the latter. For constant variability timescales, resonance is efficient for both periodic and random polarity reversals, the latter being stronger than the former. For a time variable variability timescale, resonance crossing is a more efficient excitation mechanism when polarity reversals are periodic. Each polarity reversal type has distinct  features that may help constrain the magnetic history of the star through the observation of its planetary companions.  For instance, outward planet migration to large distances from parent stars is one of the natural outcomes of periodic polarity reversal excitation if resonance crossing is sufficiently slow. Applying the excitation mechanism to the solar system, we find that the planet-jet variability resonance with periodic polarity reversal momentum loss  is a possible origin for the hitherto unexplained inclination of Jupiter's orbit by $6^\\circ$ with respect to the Sun's equator. ",
        "introduction": "The discovery of more than 760 exoplanets to date \\citep{b18} has revolutionized our understanding of the architecture of planetary systems. One of the surprising key observations is the large eccentricity of exoplanet orbits. Compared to Jupiter's eccentricity of 0.05 and to the Earth's 0.017, the median eccentricity of exoplanet orbits is 0.2. Various dynamical mechanisms were proposed to account for this departure from the solar system's planetary standard. These mechanisms include: planet-planet scattering \\citep{scatt3,scatt1,scatt2},  three-body secular Kozai oscillations \\citep{koz2,koz3,koz1}, mean motion resonances \\citep{res1,res2,res3}, stellar encounters \\citep{stenc} and excitation induced by stellar jets (\\cite{b19} hereafter Paper I, \\cite{rj8}). Although some mechanisms are more efficient than others, attention has  recently been devoted to planet-planet scattering owing to the claim that this mechanism reproduces the observed eccentricity distribution of exoplanet orbits. However, the eccentricity distribution produced from planet-planet scattering only approximates the observed distribution if the latter is cut off at an eccentricity of 0.2  \\citep{scatt2}. Smaller planetary orbital eccentricities were attributed to the other possible mechanisms or a combination thereof. The threefold increase in the number of planets between 2008 and 2012 has provided better statistics and decreased the median eccentricity from 0.3 to 0.2 (or 0.23) if only planets with periods larger than 5 (or 20) days are included to account for tidal evolution bias.   The excess of planets in the current enlarged sample with eccentricities smaller than 0.3 with respect to the Rayleigh distribution produced by planet-planet scattering \\citep{scatt2} requires that the eccentricity cutoff applied to the observed distribution be increased to 0.3 leaving more than half the planet population with stirred orbits outside the scope of applicability of planet-planet scattering.   In this paper, we revisit the excitation of planetary orbits that results from momentum loss through stellar jets (Paper I) and attempt to assess quantitatively the effects of this mechanism. Orbital excitation through stellar jets is based on jet-counterjet asymmetry  that has been observed in a significant fraction of star-disk systems \\citep{b10,b11,b12,b14,b13,hh7, rj1,rj2, rj3, rj4, rj5}. As the planet orbits around the star and the inner disk, it sees that system accelerating away from it owing to asymmetric momentum loss. Excitation from a smooth time-varying jet-induced acceleration is a secular process that requires that the jet axis be inclined with respect to the disk's plane. It was shown in Paper I that the maximum eccentricity achieved is proportional to the sine of the mutual inclination of the planetary orbital normal and the jet axis. In this paper, we develop a realistic modeling of momentum loss as a time-variable stochastic process that results in a zero mean stellar acceleration and restrict our attention to systems where the jet system's axis is perpendicular to the initial planetary plane and secular excitation is absent. Momentum loss models include periodic or random polarity reversals that may be associated with the magnetic polarity reversals of the parent star and the inner star-disk interface. In section 2, we recall the basics of orbital excitation through stellar jets. In section 3, we characterize orbital excitation as a function of the variability time scale and acceleration standard deviation as well as in the presence of mutual planetary perturbations. We identify a fundamental excitation resonance between the planet orbital period and the variability timescale and show that it is an efficient excitation mechanism for both periodic and random reversal jet profiles at constant variability timescale.  In section 4, we examine resonance crossing by modeling the time dependence of the variability time scale. In particular, it is found that resonance crossing  is an efficient excitation mechanism for periodic polarity reversal profiles but not for random polarity reversal profiles.  In the solar system, we find that resonance crossing with periodic polarity reversal and a time variability timescale that increases from 0.5 to 11 years (the current solar cycle half period) is able to reproduce Jupiter's and Saturn's orbital configuration and their inclination by $6^\\circ$ with respect to the solar equator. Section 5 contains concluding remarks. ",
        "conclusions": "In this work we examined quantitatively the excitation of planetary orbits by stellar jet stochastic momentum loss that on average does not accelerate the star. We modeled momentum loss using two main parameters, the acceleration standard deviation and the variability timescale, along with two polarity reversal modes, random and periodic. In particular, we did not invoke a prior inclination of the jet axis as it was taken to be perpendicular to the initial planetary orbits. Whereas secular excitation by asymmetric momentum loss  requires such inclination (Paper I), stochastic momentum loss does not and may achieve far greater amplitudes than secular excitation. Stochastic momentum loss  is efficient at the resonance of the planet's period with the variability timescale. Random polarity reversal appears to cause greater excitation for constant variability timescales but it fails to compete with periodic polarity reversal when the variability timescale is time dependent. If polarity reversal is related to the magnetic field of the star-disk interface, then the reversal timescale will increase during the braking of the star's rotation as indicated by observations of solar-type stars and numerical simulations of young stars' magnetic fields. As the variability timescale increases resonance crossing by the planets' orbits may excite them significantly. We have characterized such excitation and showed that the greater diversity of orbital outcomes occurs with periodic polarity reversal and is determined by how fast the planet-jet variability resonance is crossed. The smallest crossing velocities produce the most extreme systems. In particular, planets can migrate  several hundred AU away from the planet forming region around a solar mass star. Periodic polarity reversal stochastic momentum loss can statistically explain the current configuration of the solar system's Jupiter and Saturn and particularly the hitherto unknown origin of the inclination of Jupiter's orbit with respect to the solar equator. Although our study was focused on planetary systems, stellar companions of a jet-sustaining star are affected similarly by stochastic momentum loss. Perhaps the most the promising result in this work is the possible link between stellar magnetic cycles and the dynamical architecture of planetary companions. This may prove a valuable tool to constrain the magnetic history of planet-hosting or binary stars and understand the observed diversity of planetary systems."
    },
    "1208/1208.0602_arXiv.txt": {
        "abstract": "We present improved methods for using stars found in astronomical exposures to calibrate both star and galaxy colors as well as to adjust the instrument flat field. By developing a spectroscopic model for the SDSS stellar locus in color-color space, synthesizing an expected stellar locus, and simultaneously solving for all unknown zeropoints when fitting to the instrumental locus, we increase the calibration accuracy of stellar locus matching. We also use a new combined technique to estimate improved flat-field models for the Subaru SuprimeCam camera, forming `star-flats' based on the magnitudes of stars observed in multiple positions or through comparison with available measurements in the SDSS catalog. These techniques yield galaxy magnitudes with reliable color calibration ($\\lesssim 0.01$ - 0.02 mag accuracy) that enable us to estimate photometric redshift probability distributions without spectroscopic training samples. We test the accuracy of our photometric redshifts using spectroscopic redshifts $z_s$ for $\\sim$5000 galaxies in \\speczfields~cluster fields with at least five bands of photometry, as well as galaxies in the COSMOS field, finding \\mbox{$\\sigma((z_p - z_s)/(1+z_s)) \\approx $~\\accuracy} for the most probable redshift $z_p$. We show that the full posterior probability distributions for the redshifts of galaxies with five-band photometry exhibit good agreement with redshifts estimated from thirty-band photometry in the COSMOS field. The growth of shear with increasing distance behind each galaxy cluster shows the expected redshift-distance relation for a flat $\\Lambda$-CDM cosmology. Photometric redshifts and calibrated colors are used in subsequent papers to measure the masses of \\clusterfields~galaxy clusters from their weak gravitational shear and determine improved cosmological constraints. We make our Python code for stellar locus matching publicly available at \\codeurl; the code requires only input catalogs and filter transmission functions. ",
        "introduction": "\\label{sect:intro}   A principal challenge for current and planned optical and near-IR wide-field surveys is to estimate accurate redshift probability distributions for millions of galaxies from broadband photometry. Correct probability distributions are necessary, for example, for the weak lensing cosmological measurements that current and upcoming surveys (e.g., Dark Energy Survey; Large Synoptic Survey Telescope) aim to extract from wide-field optical imaging. Photometric redshift algorithms, however, can show significant systematic biases if the input galaxy photometry has even modest ($\\sim$0.03-0.04 mag) calibration error. To infer the weak lensing masses of galaxy clusters using photometric redshifts estimated from Subaru and CFHT photometry, we have developed and applied several techniques to calibrate broadband galaxy colors to an accuracy of $\\sim$0.01-0.02 magnitudes, without requiring specific standard star observations.      The relative distribution of counts recorded during flat-field exposures of an illuminated screen or of the sky can differ from the actual instrument sensitivity by up to $\\sim$10\\% across the focal plane (\\citealt{manfroid01}; \\citealt{koc03}; \\citealt{magnier04}; \\citealt{caa07}; \\citealt{rcg09}) because of a combination of geometric distortion, superposed reflections, and imperfect flat-field sources. We combine two methods to measure the Subaru SuprimeCam `star flat,' the map of the spatially dependent zeropoint error that remains after traditional flat-field correction, across a decade of observations and camera upgrades. Our star-flat model is informed by the magnitudes of the same stars observed in multiple locations on the focal plane, as well as by comparisons with available SDSS catalog magnitudes.    A powerful color calibration strategy suitable for use with medium-to-wide field data uses the fact that almost all of the stars observed in any field lie along a well-understood one-dimensional locus in color-color space (e.g., \\citealt{high09}; \\citealt{macdonald04}). According to this technique, the zeropoints of the filters are shifted until the position of the observed stellar locus matches the expected locus. The resulting calibration automatically corrects for Milky Way dust extinction. We improve the accuracy of this technique by constructing a spectroscopic model for the SDSS stellar locus, and by developing a numerical algorithm that fits for all unknown zeropoints simultaneously.   Comparison of photometric redshifts estimated from our calibrated galaxy magnitudes against spectroscopic redshifts show that these bootstrapped calibration techniques are effective. We find excellent agreement between the most probable photometric redshift $z_p$ and the spectroscopic redshift $z_s$, with a measurement error of \\mbox{$\\sigma((z_p - z_s)/(1+z_s)) \\approx $~\\accuracy}. Five-band $p(z)$ distributions summed over different sets of galaxies ($\\sum_{\\mathrm{gal}} p(z)$) show congruence with the \\citet{ics09} 30-band photometric redshift distributions for the same sets of galaxies. The growth of lensing shear with increasing redshift of galaxies behind each cluster, sensitive to the photometric redshifts of galaxies too faint to be represented in spectroscopic samples, exhibits the shape expected for a flat $\\Lambda$-CDM cosmology.    The algorithms and techniques described in this paper may be useful as a primary means of calibration or as a demanding test of zeropoint accuracy. This paper is the second in a series (``Weighing the Giants\") addressing the specific task of measuring accurate galaxy cluster masses using shear-based weak lensing methods. Paper I \\citep{vaa12} in this series describes the overall project strategy, the cluster sample and the data reduction methods. Paper III \\citep{avk12} presents a Bayesian approach to measuring galaxy-cluster masses, that uses the full photometric redshift probability distributions reported here; these masses are compared to those measured with a standard `color-cut' method based on three-filter photometry for each field.  Section \\ref{sec:sample} of this paper summarizes the wide-field imaging data used here. In Section~\\ref{sec:flat}, we describe how we determine the SuprimeCam star flats, which we use to extract consistent magnitudes across the CCD array. Section~\\ref{sec:slr} describes the stellar locus calibration algorithm and the spectroscopic model we have developed for the stellar locus. In Section~\\ref{sec:photoz}, we discuss the algorithms and the templates for galaxy spectra that we use to estimate photometric redshift probability distributions $p(z)$. A method for finding the zeropoints of  $u^*$- and $B_J$- band photometry is presented in Section~\\ref{sec:extrapolate}. In Section~\\ref{sec:confirm}, we use  the galaxy cluster red sequence and spectroscopic redshift measurements in the cluster fields to evaluate the accuracy of our photometric calibrations and redshift estimates. In Section~\\ref{sec:prob}, we compare the redshift probability distributions determined from calibrated photometry in five bands ($B_JV_Jr^+i^+z^+$) against both the zCOSMOS spectroscopic redshift sample and the most probable redshift inferred from thirty imaging bands in the COSMOS field. General agreement between the observed growth of weak lensing shear with distance behind the massive clusters and the $\\Lambda$-CDM expectation is found in Section~\\ref{sec:shearratio}. In Section~\\ref{sec:conclusions}, we summarize the calibration techniques and the quality of the photometric redshift estimates. ",
        "conclusions": "\\label{sec:conclusions}  Converting the signal measured by CCD sensors mounted on a wide-field camera to calibrated AB magnitudes with better than several percent accuracy has typically required time-consuming calibration exposures as well as optimal observing conditions. When useful overlap with survey photometry is not available (e.g., limited dynamic range, color transformations from survey magnitudes), repeated standard-star observations in photometric conditions are necessary for accurate zeropoint calibration. Dedicated exposures sequences of dense stellar fields have been considered necessary to measure the position-dependent zeropoint error ($\\sim$0.04-0.06 mag for SuprimeCam) present after correcting by dome or sky-flat calibration exposures.  For the Subaru and CFHT imaging that we analyzed, dedicated calibration data were not always available and conditions were not always photometric. We have developed improved tools that use only the instrumental stellar magnitudes in exposures of science targets to produce accurate star-flat models and color calibrations. These have sufficient power to enable high-quality photometric redshift estimates without spectroscopic training sets. We use \\speczcluster~galaxy spectra in \\speczfields~cluster fields across the sky to measure an accuracy of \\mbox{$\\sigma((z_p - z_s)/(1+z_s)) \\approx $~\\accuracy}.  We achieve $\\lesssim0.01$ - 0.02 mag color calibration by matching the observed stellar locus in color-color space to a model stellar locus. We developed a spectroscopic model for the dereddened SDSS stellar locus to estimate more accurately the locus expected for SuprimeCam and MegaPrime filters. This step is important because nonlinear color transformations exist between our SuprimeCam $B_JV_JR_CI_Cz^+$ magnitudes and the SDSS $u'g'r'i'z'$ filters used to measure the stellar locus. We created a simple $\\chi^2$ GOF statistic that makes possible a solution for all unknown zeropoints simultaneously, without correlated errors. These improvements, including the correction of the SDSS stellar locus for Galactic extinction, were necessary to estimate robust and reliable photometric redshifts from our Subaru and CFHT imaging. Interested readers may find the Python code as well as a spectroscopic model for the locus at \\codeurl. Calibration requires only a catalog of stellar magnitudes and the total transmission function for each bandpass filter.  Summarizing the redshift posterior probability distribution $p(z)$ for each galaxy as a single redshift (e.g., the most probable redshift $z_p$) disregards any probability that the true galaxy redshift may have a very different value (e.g., catastrophic outliers). The $p(z)$ distribution contains more information than the most probable redshift, or other single-point estimates. We find good agreement between the stacked $\\sum_{gal} p(z)$ and the redshift distribution estimated from thirty COSMOS photometric bands by \\citet{ics09}. The $p(z)$ distributions will provide a more powerful, less biased tool for science analysis than single-point estimates for ongoing and future surveys \\surveys.  The extreme gravitational potentials of massive galaxy clusters are expected to yield a simple, strong growth in shear signal with increasing distance of the source galaxy  behind the cluster lens. The shape of the growth of shear behind a stack of 27 clusters with increasing redshift shows agreement with a fiducial $\\Lambda$-CDM cosmology. The average shear we measure for foreground galaxies is consistent with zero.  The software we make available can produce state-of-the-art photometric accuracy without dedicated observations or the need to estimate nonlinear color terms. Stellar colors will be a useful tool for checking the accuracy of zeropoints obtained through ubercalibration of surveys (e.g.,  \\citealt{ivezic07}; \\citealt{padmanabhan08}; \\citealt{schlafly12}) or close monitoring of atmospheric transparency (e.g., \\citealt{bab10}; \\citealt{blake11})."
    },
    "1208/1208.4143.txt": {
        "abstract": "We present deep {\\em HST/WFPC2}, rest-frame $U$ images of 17 $\\sim L^\\star$ quasars at \\zo\\ and \\zt\\ ($V$ and $I$ bands respectively), designed to explore the host galaxies. We fit the images with simple axisymmetric galaxy models, including a point-source, in order to separate nuclear and host-galaxy emission. We successfully model all of the host galaxies, with luminosities stable to within 0.3~mag. Combining with our earlier {\\em NICMOS} rest-frame optical study of the same sample, we provide the first rest-frame $U-V$ colours for a sample of quasar host galaxies. While the optical luminosities of their host galaxies indicate that they are drawn purely from the most massive ($\\gtsim L^\\star$) early-type galaxy population, their colours are systematically bluer than those of comparably massive galaxies at the same redshift. The host galaxies of the radio-loud quasars (RLQ) in our sample are more luminous than their radio-quiet quasar (RQQ) counterparts at each epoch, but have indistinguishable colours, confirming that the RLQ's are drawn from only the most massive galaxies ($10^{11}-10^{12}$~M$_\\odot$ even at \\zt), while the RQQ's are slightly less massive ($\\sim10^{11}$~M$_\\odot$). This is consistent with the well-known anti-correlation between radio-loudness and accretion rate. %??? Using simple stellar population ``frosting'' models %with a 10~Myr old starburst, we estimate that $\\sim1$\\% of the mass of the host galaxies is undergoing star formation at \\zt\\ and $\\sim0.1$\\% at \\zo. The we estimate mean star formation rates of $\\sim350$~M$_\\odot$~yr$^{-1}$ for the RLQ's and $\\sim100$~M$_\\odot$~yr$^{-1}$ for the RQQ's at \\zt. By \\zo, these rates have fallen to $\\sim150$~M$_\\odot$~yr$^{-1}$ for the RLQ's and $\\sim50$~M$_\\odot$~yr$^{-1}$ for the RQQ's. We conclude that while the host galaxies are extremely massive, they remain actively star-forming at, or close to, the epoch of the quasar. ",
        "introduction": "Comprehensive studies of low redshift ($z<0.4$) quasar host galaxies indicate that the most powerful nuclear activity in present-day galaxies is usually associated with massive, bulge-dominated host galaxies~\\citep[\\eg][]{disney+95,bahcall+97,hooper+97,mcleod+99,mclure+99,ridgway+01,hamilton+02,dunlop+03,floyd+04,guyon+06,canalizo+07,bennert+08}. This ties in well with the discovery of the black hole -- bulge mass relationship in quiescent galaxies: the host galaxies are exactly the type of systems in which one would expect to find the very massive black holes that are required to power luminous quasars. However, this conclusion is based largely on their luminosities and morphologies in single-band Hubble Space Telescope (HST) imaging (especially at the highest quasar luminosities--~\\citealt{floyd+04}). The stellar populations and masses of quasar host galaxies remain poorly constrained, although ground-based $K$-band imaging~\\citep{taylor+96} and off-nuclear optical spectroscopy~\\citep{nolan+01} suggest the presence (albeit at low level) of an intermediate age stellar population in some. It is difficult to disentangle any young stellar population in the host galaxy from the glare of the active galactic nucleus (AGN), and much of our picture of quasar host galaxies is inferred from the properties of their inactive, giant red elliptical, counterparts.   Meanwhile, lower luminosity AGN in the local Universe have clear evidence for younger stellar populations~\\citep[\\eg][]{borosonoke82}. Seyfert galaxies are frequently found with circumnuclear starbursts~\\citep{heckman+97,gu+01,gonzalezdelgado+01,cidfernandes+04,davies+07,riffel+09}. The presence of intermediate age stellar populations has also been extensively shown in type 2 (obscured) and lower luminosity type 1 quasars~\\citep{kotilainenward94,ronnback+96,brotherton+99,kauffmann+03,jahnke+04a,jahnke+04b,sanchez+04,vandenberk+06,jahnke+07,jahnke+09}. A decrease in mean stellar age with increasing AGN luminosity is suggested (e.g.~\\citealt{kauffmann+03,vandenberk+06}), indicating that we may be missing young stellar populations in the most luminous nearby quasars, due to the glare of the nucleus.  At higher redshift we see greater levels of star formation, as expected. \\citet{mainieri+11} examined $\\sim 1800$ obscured (X-ray selected) quasars in the XMM-COSMOS survey, with bolometric luminosities $>8\\times10^{45}$~erg~s$^{-1}$. They confirm that all dwell in massive ($>10^{10}$~M$_\\odot$) galaxies, with a monotonically increasing AGN fraction with stellar mass. They also find strong star formation in the majority, with the star forming fraction increasing strongly with redshift (62\\% star forming at $z\\sim1$, 71\\% at $z\\sim 2$, 100\\% at $z\\sim3$). \\citet{trichas+12}, however, show a decreasing fraction of measurable starburst contributions with increasing AGN luminosity.  %In radio-loud AGN, the picture is clearer, with a massive, bulge-doiminated host galaxy (and associated supermassive black hole) being a pre-requisite for such activity~\\citep[\\eg][]{dunlop+03, floyd+04, madrid+06} %??? RLAGN???  A key observable is the level of star formation present in quasar host galaxies when compared to inactive galaxies of the same mass at the same epoch. In this paper we aim to explore the level of star formation in our \\zo\\ and \\zt\\ quasar samples, originally imaged in the rest-frame optical using HST's Near Infrared Camera and Multi-Object Spectrometer (NICMOS -- see~\\citealt{kukula+01} -- hereafter K01). That paper showed the quasars to be hosted by massive ($\\sim10^{11}-10^{12}$~M$_\\odot$) elliptical galaxies, with masses consistent with those found at lower redshift (\\citealt{mclure+99, dunlop+03, floyd+04} -- hereafter M99; D03; F04 respectively) assuming passively evolving stellar populations. Here we present the follow-up rest-frame $U$-band imaging taken with WFPC2, explore the colours and morphologies of the host galaxies, and constrain their masses and young stellar populations. In particular, we explore their colours in comparison to the colours of massive galaxies at the same epoch.  The layout of the remainder of the paper is as follows. The sample design and observations are discussed in section 2. We describe the data reduction and analysis in section 3. Our results are presented in section 4, and their implications discussed in section 5. Our main conclusions are summarised in section 6. We assume throughout a cosmology with $\\Omega_{m}=0.3$, $\\Lambda= 0.7$ and $H_{0}=70$~km s$^{-1}$ Mpc$^{-1}$, and convert previous work to this standard where required.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  \\begin{figure*} \\centering {\\includegraphics[width=85mm,angle=0]{figs/f1a.eps}} \\hspace{5mm} {\\includegraphics[width=85mm,angle=0]{figs/f1b.eps}} \\caption{\\label{fig-samp}{\\bf Left:} Absolute $V$ magnitude versus redshift for the quasars in our host galaxy studies out to $z\\approx2$. Filled circles represent radio-quiet quasars, while open circles represent radio-loud quasars. Objects in the present WFPC2 study at \\zo\\ \\& 2 are shown using large symbols, compared to our lower redshift objects (F04, D03, M99). Our sample spans the knee of the quasar luminosity function at each redshift: the dotted line indicates $M_V^\\star$ for the 2QZ quasar luminosity function at each redshift~\\citep{croom+04}. See section~\\ref{sec-samp} for notes on the sample selection. {\\bf Right:} Illustration showing generic spectra for a quasar nucleus (upper panel) and an early-type galaxy (lower panel), with our approximate rest-frame bandpasses marked. We have tailored our filter selection to target the rest-frame $U$ and $V$-band, thus sampling the SED of the host on either side of the break feature at 4000\\AA. While the Mg~II quasar emission line is admitted, we avoid prominent galaxy emission lines (section~\\ref{sec-filt}).} \\end{figure*}    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\begin{table} \\begin{center} \\begin{small} \\caption{\\label{tab-samp} Optical and radio properties of the quasars in the current study. We show the target names used in the HST archive. Redshifts, $V$-band magnitudes, colour excesses $E(B-V)$~\\citep{schlegel+98} are derived from NED. Absolute magnitudes are confined to the range $-23.9 \\leq M_{V} \\leq -25.1$ (in rest-frame) to ensure that the quasars in each redshift bin are comparable with each other (as well as with the low-redshift quasars in our earlier studies, F04, D03, M99). RQQs have $L_\\mathrm{5GHz}  < 10^{24.5}$~W~Hz$^{-1}$~sr$^{-1}$. RLQs have $L_\\mathrm{5GHz} > 10^{24.5}$~W~Hz$^{-1}$~sr$^{-1}$ and steep radio spectra. For the estimation of luminosities at both radio and optical wavelengths we have assumed a quasar spectrum of the form $F_\\nu\\propto\\nu^{-0.5}$.} See section~\\ref{sec-samp} for full details of sample selection. %      {\\bf Redshift references:} %      $^{\\mathrm (a)}$~\\cite{HB_QSO_89}; %      $^{\\mathrm (b)}$~\\cite{SDSS_QSO_DR7}; %      $^{\\mathrm (c)}$~\\cite{VCV93}; %      $^{\\mathrm (d)}$~\\cite{Croom:2dfQSO}; %      $^{\\mathrm (e)}$~\\cite{Boyle:UVQSO}; %      $^{\\mathrm (f)}$~\\cite{Dunlop:PKS_QSO}; %      $^{\\mathrm (g)}$~\\cite{Stickel:radio}. \\centering \\begin{tabular}{lccrrr} \\hline \\hline Object \t& Type \t&$z$\t&$V$ \t&$E(B-V)$\t&$S_\\mathrm{1.4}$\\\\ &      \t\t&     \t& /mag\t&  / mag\t\t&/mJy\\\\ \\hline \\multicolumn{6}{c}{\\bf \\zo\\ sample (F606W `Wide $V$')}\\\\ BVF225        \t\t& RQQ \t& 0.910 & 19.6 & 0.012&$<0.15$ \\\\ BVF247        \t\t& RQQ \t& 0.890 & 19.2 & 0.014&$<0.15$ \\\\ BVF262        \t\t& RQQ \t& 0.970 & 19.6 & 0.008&$<0.15$ \\\\ PKS0440$-$00\t\t& RLQ \t& 0.844 & 19.2 & 0.053& 1443\\\\ PKS0938$+$18\t& RLQ \t& 0.940 & 19.1 & 0.030& 432\\\\ 3C422         \t\t& RLQ \t& 0.942 & 19.5 & 0.055& 2117\\\\ MC2112$+$172\t& RLQ \t& 0.878 & 18.7 & 0.134& 425\\\\ 4C02.54\t\t\t& RLQ \t& 0.976 & 19.0 & 0.045& 790\\\\ \\multicolumn{6}{c}{\\bf \\zt\\ sample (F814W `Wide $I$')}\\\\ SGP2:36\t\t\t& RQQ \t& 1.773 & 20.8 & 0.014&$<0.45$ \\\\ SGP2:25\t\t\t& RQQ \t& 1.869 & 20.8 & 0.016&$<0.45$ \\\\ SGP2:11\t\t\t& RQQ \t& 1.976 & 21.1 & 0.017&$<0.45$ \\\\ SGP3:39\t\t\t& RQQ \t& 1.959 & 20.9 & 0.028&$<0.45$ \\\\ SGP4:39\t\t\t& RQQ \t& 1.721 & 20.9 & 0.022&$<0.45$ \\\\ PKS1524$-$13\t\t& RLQ \t& 1.687 & 21.0 & 0.121& 2400\\\\ B2~2156+29\t\t& RLQ \t& 1.753 & 17.5 & 0.104& 1223\\\\ PKS2204$-$20\t\t& RLQ \t& 1.923 & 20.3 & 0.049& 700\\\\ 4C45.51\t\t\t& RLQ \t& 1.992 & 20.6 & 0.122& 1836\\\\ \\hline \\end{tabular} \\end{small} \\end{center} \\end{table} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec-conc} The main result of this study is that UV-luminous host galaxies are present in our entire sample of radio-loud and radio-quiet quasars at \\zo\\ and \\zt. Our fitted host-galaxy luminosities are stable to within $\\pm0.3$~mag. and our conclusions are robust to changes in the preferred UV host-galaxy size. They occupy elliptical galaxies in the upper regions of the optical luminosity function ($1-4L_\\star$), but have bluer than median colours for galaxies of the same redshift and luminosity. They are strongly star forming at \\zt\\ ($\\sim 1\\%$ by mass), with weaker ($\\sim0.1\\%$) star-formation in somewhat more massive galaxies by \\zo. The RLQ's and RQQ's are indistinguishable in terms of colour, but the RLQ's are more massive at each epoch, indicative of their lower accretion rates (see section~\\ref{sec-RL}). % (although we note the possible bias in the RLQ/RQQ divide at \\zt~--- see sections~\\ref{sec-presamp, sec-kc, sec-RL}). The most massive RLQ's have stellar masses of $\\sim10^{12}$~M$_\\odot$, even at \\zt. % and are accreting at lower Eddington rates than their RQQ counterparts.  %\\citet{sanchez+04, jahnke+04b, jahnke+09}??? % %Accretion rate / radio-loudness: \\citet{panessa+07, sikora07, xu+99}  The probability of obtaining galaxies as UV-bright as these quasar hosts from a sample of quiescent galaxies of the same optical luminosity and redshift is low ($<0.12\\%$), and we reject the hypothesis that quasar host galaxies are drawn at random from the massive galaxy population. We conclude that while the host galaxies are extremely massive, they remain actively star-forming (or very recently so) at the epoch of the quasar. The relationship between AGN activity and star-formation remains unknown, but it seems clear that optically luminous AGN activity is associated with recent or ongoing star-formation in massive galaxies, whereas radio AGN activity is associated purely with the most massive galaxies.  The host galaxies are more compact than low-redshift analogues, but are consistent with massive quiescent galaxies at same epoch. There is clear evidence for clumpiness in the UV flux. We reiterate that the models are designed to fit the smooth bulk of the starlight. However, we have noted that the radial profiles of the host galaxies presented here are not as smooth as those in the optical. The UV ``bumps'' correspond to compact regions of increased UV luminosity -- perhaps close companion objects or knots of star formation in the host galaxy. These were not apparent in the NICMOS rest-frame V imaging. The UV clumpiness, and compact UV morphology observed in some objects suggest that higher resolution imaging and IFU spectroscopy may reveal compact regions of star formation in the host galaxy, and hint at the possibility of discovering circumnuclear starbursts in some particularly UV-compact objects. Our results therefore support the growing consensus that quasar and AGN activity is strongly associated with recent star formation (possibly its truncation -- see section~\\ref{sec-RL}), and suggest that this activity tends to be located in the central regions of the galaxy.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1208/1208.5482_arXiv.txt": {
        "abstract": "Dark halo substructure may reveal itself through secondary, small-scale gravitational lensing effects on light sources that are macrolensed by a foreground galaxy. Here, we explore the prospects of using Very Long Baseline Interferometry (VLBI) observations of multiply-imaged quasar jets to search for submilliarcsecond-scale image distortions produced by various forms of dark substructures in the $10^3$--$10^{8}\\ M_\\odot$ mass range. We present lensing simulations relevant for the angular resolutions attainable with the existing European VLBI Network (EVN), the global VLBI array, and an upcoming observing mode in which the Atacama Large Millimeter Array (ALMA) is connected to the global VLBI array. While observations of this type would not be sensitive to standard cold dark matter subhalos, they can be used to detect the more compact forms of halo substructure predicted in alternative structure formation scenarios. By mapping $\\approx 5$ strongly lensed systems, it should be possible to detect or robustly rule out primordial black holes in the $10^3$--$10^6\\ M_\\odot$ mass range if they constitute $\\gtrsim 1\\%$ of the dark matter in these lenses. Ultracompact minihalos are harder to detect using this technique, but $10^6$--$10^8\\ M_\\odot$ ultracompact minihalos could in principle be detected if they constitute $\\gtrsim 10\\%$ of the dark matter. ",
        "introduction": "\\label{intro} A generic prediction of the standard cold dark matter (CDM) scenario is that a substantial fraction of the total mass of galaxy-sized dark matter halos \\citep[$\\sim 10\\%$; ][]{Gao et al.,Maciejewski et al.} should be in the form of bound substructures (a.k.a. subhalos or subclumps) left over from the process of hierarchical assembly. The fact that the number of substructures seen in CDM simulations greatly outnumber the satellite galaxies detected in the vicinity of the Milky Way and Andromeda constitutes the so-called ``missing satellite problem'' \\citep{Klypin et al.,Moore et al.}. While it has been argued that astrophysical processes that quench star formation in low-mass halos may explain this discrepancy \\citep[e.g.][]{Maccio et al. b,Font et al.}, this implies that a vast population of extremely faint or completely dark substructures should be awaiting discovery in the halos of galaxies. Provided that CDM is in the form of Weakly Interacting Massive Particles (WIMPs), these subhalos are in principle detectable by the Fermi Gamma-ray Space Telescope because of their annihilation fluxes. However, Fermi has so far failed to detect any unambigious signal from such objects \\citep[e.g.][ but see \\citealt{Bringmann et al. b} and \\citealt{Su & Finkbeiner} for a different view]{Belikov et al.,Zechlin et al.,Hooper & Linden}.  Gravitational lensing may provide an independent test for the presence of dark halo substructures \\citep[for a review, see][]{Zackrisson & Riehm}. A foreground galaxy that happens to be aligned with a background light source can produce multiple images of the background object, with a typical image separation of $\\sim 1\\arcsec$ (an effect known as strong lensing or macrolensing). While simple, smooth models of galaxy lenses are usually able to reproduce the positions of these macroimages, their observed flux ratios are more difficult to explain. Such flux-ratio violations have been interpreted as evidence of substantial small-scale structure within the main lens \\citep[e.g.][]{Mao & Schneider,Chiba,Keeton et al.,Kochanek & Dalal}. A notable problem with this picture is that current CDM simulations predict too little substructure to explain many of these flux-ratio violations (e.g. \\citealt{Maccio & Miranda,Xu et al. a,Chen et al.} -- but see \\citealt{Metcalf & Amara}), possibly pointing to a considerable contribution from low-mass halos elsewhere along the line of sight \\citep{Xu et al. b} or some additional form of substructure (dark or luminous) within the lens.  A slightly different lensing approach exploits the small-scale distortions that halo substructure is expected to introduce in the morphologies of extended macroimages.  Substructures of mass $\\gtrsim 10^8\\ M_\\odot$ can perturb gravitational arcs and Einstein rings on scales resolvable with the {\\it Hubble Space Telescope} \\citep[][]{Vegetti & Koopmans a,Vegetti & Koopmans b} and  detections of $\\sim 10^9$--$10^{10}\\ M_\\odot$ objects have already been made this way \\citep{Vegetti et al. a,Vegetti et al. b,Vegetti et al. c}. In line with the flux ratio anomaly results, these observations seem to suggest a subhalo mass fraction that is significantly higher than predicted by standard CDM, and possibly also a flatter subhalo mass function slope \\citep{Vegetti et al. b, Vegetti et al. c}.  By mapping extended macrolensed sources with milliarcsecond or sub-milliarcsecond resolution using Very Long Baseline Interferometry (VLBI) techniques at radio wavelengths, substructures at even lower masses can in principle be detected. Such objects may introduce kinks and bends in multiply-imaged quasar jets \\citep{Wambsganss & Paczynski,Metcalf & Madau} and one detection of a $\\sim 10^5$--$10^7 M_\\odot$ object has already been claimed using this technique \\citep{Metcalf}. In this situation, the lensing effects produced by halo substructures can be separated from intrinsic morphological features in jets, since the latter would be reproduced in all macroimages whereas dark matter clumps in the halo of the lens would affect each macroimage differently. Similar methods for exploiting the lensing effects produced by halo substructures on scales of $\\sim 100$ milliarcseconds down to $\\sim 0.01$ milliarcseconds have also been explored by \\citet{Yonehara et al.,Inoue & Chiba a,Inoue & Chiba b,Inoue & Chiba c, Hisano et al.,Ohashi et al.,Riehm et al.} and \\citet{Hezaveh et al.}. However, effects of this type tend to be sensitive to the density profiles of substructures, and may be undetectable for all but the very densest, most extreme forms of substructure \\citep{Zackrisson et al.}.  Here, we use lensing simulations to explore the prospects of using macrolensed quasar jets observed at sub-milliarcsecond resolution, in searches for standard CDM subhalos, ultracompact minihalos and primordial black holes within the main lens. These different forms of substructure are described, along with previous constraints on such objects, in Sect.~\\ref{substructures}. The details of our simulations and assumptions are covered in Sect.~\\ref{simulations}. In Sect.~\\ref{results}, we present our results and in Sect.~\\ref{discussion} we discuss some lingering issues with the adopted technique. Sect.~\\ref{summary} summarizes our findings. ",
        "conclusions": "\\label{discussion}  \\subsection{Temporal effects} In previous sections, we have argued that millilensing-induced distortions of quasar jets may be distinguished from morphological features intrinsic to these sources, since the latter would be reproduced in all macroimages whereas millilensing should affect each image independently. However, this argument comes with a caveat. The time delay between the images in quasar-galaxy lenses can be up to a year \\citep[for a compilation of time delays, see][]{Oguri}, which means that intrinsic, transient features in the jet may, at any given time, be visible in just one of the images and be mistaken for millilensing effects. This is for instance likely to be the case in superluminal radio jets, where blobs are seen to move $\\sim 1$ milliarcseconds yr$^{-1}$ along the jet \\citep[e.g.][]{Jorstad et al.}. For macrolensed jets that show signs of millilensing distortions, it may therefore become necessary to obtain data at two or more epochs. Since halo substructures give rise to millilensing magnification pattern that will appear stationary over decades \\citep{Metcalf & Madau}, any distortions that seem to move along the jet are bound to be intrinsic to the source. Small-scale features that are not duplicated in the other macroimages and appear with a fixed angular position (as, for instance, measured from the base of the jet) over the course of more than a year is on the other hand likely caused by millilensing. \\begin{table} \\caption{Impact parameter $R_\\mathrm{eff}$ within which a subhalo of a given type will produce detectable macroimage distortions} \\begin{tabular}{@{}lllll@{}} \\hline Type & Resolution    & Mass        & $\\mu$ & $R_\\mathrm{eff}$ \\\\ & (milliarcsec) & ($M_\\odot$) &       & (pc)\\\\ \\hline IMBH & 0.05 (86 GHz) & $10^3$ & 3  & 1\\\\ & \t\t\t\t\t\t\t & $10^4$ &    & 3\\\\ & \t\t\t\t\t\t\t & $10^5$ &    & 7\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 20\\\\ &               & $10^3$ & 10 & 2\\\\ & \t\t\t\t\t\t\t & $10^4$ &    & 6\\\\ & \t\t\t\t\t\t\t & $10^5$ &    & 20\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 50\\\\ &               & $10^3$ & 30 & 2\\\\ & \t\t\t\t\t\t\t & $10^4$ &    & 10\\\\ & \t\t\t\t\t\t\t & $10^5$ &    & 40\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 80\\\\ & 0.3 (22 GHz)\t & $10^4$ & 3  & 2\\\\ & \t\t\t\t\t\t\t & $10^5$ &    & 5\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 20\\\\ &    \t \t\t\t\t & $10^4$ & 10 & 2\\\\ & \t\t\t\t\t\t\t & $10^5$ &    & 7\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 30\\\\ &    \t \t\t\t\t & $10^4$ & 30 & 2\\\\ & \t\t\t\t\t\t\t & $10^5$ &    & 8\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 40\\\\ & 0.7 (8.4 GHz) & $10^5$ & 3  &  3\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 10\\\\ &    \t \t\t\t\t & $10^5$ & 10 &  4\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 20\\\\ &    \t \t\t\t\t & $10^5$ & 30 &  8\\\\ & \t\t\t\t\t\t\t & $10^6$ &    & 40\\\\ UCMH & 0.05 (86 GHz) & $10^6$ & 3  &  2\\\\ & \t\t\t\t\t\t\t & $10^7$ &    &  6\\\\ & \t\t\t\t\t\t\t & $10^8$ &    &  20\\\\ &               & $10^6$ & 10 &  3\\\\ & \t\t\t\t\t\t\t & $10^7$ &    &  20\\\\ & \t\t\t\t\t\t\t & $10^8$ &    &  60\\\\ &               & $10^6$ & 30 &  10\\\\ & \t\t\t\t\t\t\t & $10^7$ &    &  30\\\\ & \t\t\t\t\t\t\t & $10^8$ &    &  100\\\\ & 0.3 (22 GHz)\t & $10^7$ & 3  &  2\\\\ & \t\t\t\t\t\t\t & $10^8$ &    &  10\\\\ &    \t \t\t\t\t & $10^7$ & 10 &  4\\\\ & \t\t\t\t\t\t\t & $10^8$ &    &  30\\\\ &    \t \t\t\t\t & $10^7$ & 30 &  5\\\\ & \t\t\t\t\t\t\t & $10^8$ &    &  60\\\\ & 0.7 (8.4 GHz) & $10^8$ & 3  &  10\\\\ &    \t \t\t\t\t & $10^8$ & 10 &  20\\\\ &    \t \t\t\t\t & $10^8$ & 30 &  60\\\\ \\hline \\label{Reff} \\end{tabular} \\end{table}  \\subsection{Source size sensitivity} \\label{size} For a fixed substructure type and telescope beam size, the prospects of detecting millilensing effects depend on the adopted source dimensions. This is exemplified in Fig.~\\ref{size_example}, where the probability of substructures is seen to increase with source area -- whereas only one IMBH is detectable in small-source case (left), two IMBHs are detectable for the larger source (right). The source sizes adopted in Sect.~\\ref{source size} are uncertain by a factor of a few, and it may be convenient to be able to generalize our results to match other source dimensions. For a given substructure type and mass, any detection limit $\\min(f_\\mathrm{sub, 1})$ (as listed in Table ~\\ref{IMBH_limits} and ~\\ref{UCMH_limits}) derived for an intrinsically elliptical source with area $A_1$ can be rescaled to some other source area  $A_2$ using: \\begin{equation} \\min f_\\mathrm{sub, 2} \\approx \\frac{A_1 + C_1 R_\\mathrm{eff}}{A_2 + C_2 R_\\mathrm{eff}} \\min f_\\mathrm{sub, 1} \\label{rescaling_formula} \\end{equation} Here, $\\min f_\\mathrm{sub, 2}$ is the rescaled detection limit relevant for source area $A_2$, whereas the $C$ parameters represent the circumferences of the macorimages one wants to rescale from ($C_1$) and to ($C_2$). The impact parameter  $R_\\mathrm{eff}$ measures the  projected distance from the subhalo centre within which detectable macroimage distortions will be produced. This impact parameter, which depends on both subhalo mass and type, is typically larger than the subhalo Einstein radius, since substantial deflection can occur even outside the latter. The $R_\\mathrm{eff}$ values relevant for $10^3$--$10^6\\ M_\\odot$ IMBHs and $10^6$--$10^8\\ M_\\odot$ UCMHs are listed in Table~\\ref{Reff} for the resolutions adopted at 8.4, 22 and 86 GHz. Since these $R_\\mathrm{eff}$ values also depend on the magnification of the macroimage, $R_\\mathrm{eff}$ values are presented for $\\mu=3$, 10 (our default value) and 30.  This rescaling scheme, which assumes that the source size and $R_\\mathrm{eff}$ are independent, is admittedly an approximation and reliable only to within a factor of a few. Secondary images due to substructure lensing may for instance be easier to detect for a compact rather than an extended source due to flux ratio issues. An effect of the latter type is evident in Table~\\ref{Reff}, where both $10^6$ IMBHs and $10^8$ UCMH are seen to have larger $R_\\mathrm{eff}$ at 86 GHz (smallest source) than at 22 (intermediate source) or 8.4 GHz (largest source).  \\begin{figure} \\includegraphics[scale=0.55]{f7.eps} \\caption{Illustration of how source size affects the probability for detecting dark halo substructure. The two frames depict a single macroimage (out of a two-image pair) with fixed macrolensing magnification ($\\mu=10$) but with different intrinsic source size: $10 \\times 2.5$ pc (left) and $40 \\times 10$ pc (right). The smaller version (left) corresponds to the source size adopted for our 22 GHz simulations (see Sect.~\\ref{source size}). The red dots mark the positions of two $10^5 \\ M_\\odot$ IMBHs (identical positions in both frames). In the small-source case (left), only one of these IMBHs produce detectable millilensing effects, whereas both can be detected in the large-source case (right) due to better macroimage coverage of the lens plane. A resolution of 0.3 milliarcsec has been adopted in both cases (as considered suitable for 22 GHz observations). \\label{size_example}} \\end{figure}  \\begin{figure*} \\includegraphics[scale=0.55]{f8a.eps} \\includegraphics[scale=0.55]{f8b.eps} \\caption{Examples of how source morphology and surface brightness distribution affect the detectability of millilensing from IMBHs of mass $10^5\\ M_\\odot$ and $10^6\\ M_\\odot$ against a simulated radio map of a strongly lensed quasar jets at 8.4 GHz. For each IMBH mass, the leftmost frame contains a source with constant surface brightness, the middle one the Gaussian profile used throughout the rest of the paper, and the rightmost frame a source consisting of a sequence of Gaussian ``blobs\" of different sizes. Contour representations are not used in this plot, since this becomes confusing in the case of a flat surface brightness profile. In general, complicated source morphologies (rightmost frames) do not significantly compromise the detectability of millilensing effects, but sources with shallow (or even constant) surface brightness profiles may render certain forms of substructure lenses undetectable (as seen in the case where $10^5\\ M_\\odot$ IMBH are superposed on a constant surface brightness source).} \\label{sfb} \\end{figure*}  \\subsection{The surface brightness profile} \\label{sbp} In previous sections, we have assumed the source to be an intrinsically straight jet with a surface brightness distribution described by a 2-dimensional Gaussian. While the intrinsic source morphology and surface brightness profile is less important than the overall source area when assessing lensing probabilities, there are certain situations where they do matter. Since gravitational lensing conserves surface brightness, halo substructure can only produce detectable image distortions if it happens to affect a region of the macroimage where there is a non-negligble surface brightness gradient. In the extreme case of a source with constant surface brightness, halo substructure will not produce {\\it any} detectable features unless its lensing effects extends beyond the macroimage boundary. This is exemplified in Fig.~\\ref{sfb}, where $10^5\\ M_\\odot$ and $10^6\\ M_\\odot$ IMBHs (red dots) are superposed on macroimages of (from left to right, for each IMBH mass) an elliptical source with constant surface brightness, an elliptical source with Gaussian surface brightness profile and a patchy jet with Gaussian ``blobs\" of increasing size when moving from the lower-right to upper-left corner.  In this example, a $10^5\\ M_\\odot$ IMBH becomes undetectable in the case of a constant surface brightness source, but can be spotted as a mild distortion against the Gaussian source. An IMBH of this mass redistributes surface brightness within an area that is much smaller than that of the source. Hence, if the source surface brightness is constant, no detectable effects are produced. Even though placed in the exact same position, the lensing produced by a $10^6\\ M_\\odot$ IMBH on the other hand extends sufficiently far out to distort the rim of the macroimage and can therefore be detected regardless of the source profile. In fact, the only $f_\\mathrm{IMBH}$ entry in Table ~\\ref{IMBH_limits} that would change in any dramatic way when going from a Gaussian to a constant surface brightness source corresponds to the $10^5\\ M_\\odot$ case depicted in Fig.~\\ref{sfb} (i.e. source size and resolution corresponding to 8.4 GHz). In this case, constant surface brightness source would effectively prevent any useful $f_\\mathrm{IMBH}$ constraints, whereas the changes are modest in all other cases. Since UCMH lenses produce more long-range effects then IMBHs, the $f_\\mathrm{UCMH}$ estimates in Table ~\\ref{UCMH_limits} are even less affected by the source surface brightness profile.  Fig.~\\ref{sfb} also provides an example of a more patchy jet morphology. This jet has the same source area as the other source cases, and consequently extends further in the vertical direction due to the empty regions between the ``blobs\". Both $10^5\\ M_\\odot$ and  $10^6\\ M_\\odot$ IMBHs are in principle detectable against the source in this example, although the distortion produced in the former case becomes very modest since the IMBH happens to be projected on the outskirts of one of the blobs. In general, having a complicated jet morphology does not significantly compromise the detectability of millilensing effects. Instead, a morphology of this type could even boost the detection prospects in cases where the substructure $R_\\mathrm{eff}$ (see Sect.~\\ref{size}) is larger than empty regions in the macroimage (as in the $10^6\\ M_\\odot$ IMBH case in Fig.~\\ref{sfb}), since the effective source area becomes larger in this situation.   \\subsection{The nature of the substructures} The detection of milliarcsecond or submilliarcsecond-scale image distortions would prove the existence of substructures within the macrolens, and also allow constraints on their surface number densities (as a function of substructure mass and type) to be set. However, the exact nature of a single millilens may still be very difficult to determine, since a low-mass, high-density substructure can produce a distortion very similar to that of a high-mass, low-density object. While \\citet{Inoue & Chiba c} have demonstrated that the distortions induced in extended images (like the ones we model here) contain some information about the density profiles of the lenses, the finite resolution and sensitivity of actual observations could still allow for considerable degeneracies in cases where neither the masses nor the density profiles of the millilenses are known. IMBHs and UCMHs can for instance produce very similar lensing distortions in our simulations (although at different masses -- a UCMH typically needs to be a factor of $\\sim 10^3$ more massive than an IMBH to reproduce a given feature). While it is possible that a combined consideration of small-scale distortions (e.g. the bending of a macrolensed jets), astrometric perturbations (the positional shift of a macroimage produced by the presence of substructures) and macroimage flux ratios could provide some constraints, this is beyond the scope of the present paper."
    },
    "1208/1208.2094_arXiv.txt": {
        "abstract": "We perform multi-dimensional, time-dependent radiation transfer simulations for hard X-ray and $\\gamma$-ray emissions, following radioactive decays of $^{56}$Ni and $^{56}$Co, for two-dimensional delayed detonation models of Type Ia supernovae (SNe~Ia). The synthetic spectra and light curves are compared with the sensitivities of current and future observatories for an exposure time of $10^6$ seconds. The non-detection of the $\\gamma$-ray signal from SN 2011fe at 6.4 Mpc by SPI on board {\\em INTEGRAL} places an upper limit for the mass of $^{56}$Ni of $\\lesssim 1.0$ \\msun, independently from observations in any other wavelengths. Signals from the newly formed radioactive species have not been convincingly measured yet from any SN~Ia, but the future X-ray and $\\gamma$-ray missions are expected to deepen the observable horizon to provide the high energy emission data for a significant SN~Ia sample.  We predict that the hard X-ray detectors on board {\\em NuStar} (launched in 2012) or {\\em ASTRO-H} (scheduled for launch in 2014) will reach to SNe~Ia at $\\sim$15 Mpc, i.e., one SN every few years. Furthermore, according to the present results, the soft $\\gamma$-ray detector on board {\\em ASTRO-H} will be able to detect the 158 keV line emission up to $\\sim$25 Mpc, i.e., a few SNe Ia per year.  Proposed next generation $\\gamma$-ray missions, e.g., {\\em GRIPS}, could reach to SNe Ia at $\\sim$$20 - 35$ Mpc by MeV observations. Those would provide new diagnostics and strong constraints on explosion models, detecting rather directly the main energy source of supernova light. ",
        "introduction": "It is widely accepted that Type Ia Supernovae (SNe Ia) are a major source of Fe in the Universe \\citep[see e.g.,][for a review]{hillebrandt2000}.  The thermonuclear explosion of a white dwarf produces $^{56}$Ni as a main product \\citep[e.g.,][]{nomoto1984}.  It decays into $^{56}$Co (with an e-folding time of $\\sim$8.8 days) and then into $^{56}$Fe ($\\sim$113 days). The transitions typically occur into excited states of the daughter nuclei, which generally de-excite by emissions of $\\gamma$-rays with characteristic energies of $\\sim$1 MeV. These $\\gamma$-rays energize the thermal electron pool mainly through Compton scattering, which ultimately leads to the optical appearance of SNe Ia.  Although this scenario has been supported by studying optical spectra and light curves, the most direct evidence is still missing, i.e., the detection of the decay $\\gamma$-rays and related high energy emissions \\citep[e.g.,][]{clayton1969,ambwani1988,milne2004}.  A marginal detection for the peculiar SN Ia 1991T was reported \\citep{lichti1994,morris1997}, while for the more-nearby SN Ia 1998bu only upper limits were obtained \\citep{georgii2002}.  The recently discovered SN Ia 2011fe in the nearby galaxy M101 at $\\sim$6.4 Mpc \\citep{nugent2011a,nugent2011b} has been observed by {\\em INTEGRAL}, but only an upper limit has been placed \\citep{isern2011a,isern2011b}.  Despite the importance of multi-dimensional structures in state-of-the-art explosion models \\citep[e.g.,][]{gamezo2003, roepke2005,bravo2006,roepke2007,jordan2008,seitenzahl2011} up to now most theoretical studies of the high energy emission of SNe~Ia have been restricted to one-dimensional models \\citep[e.g.,][]{ambwani1988,hoflich1992, hoflich1998,gomez1998}. For a review of the (one-dimensional) theoretical studies of high energy signals from SNe~Ia see \\citet{milne2004}. Only recently, first multi-dimensional studies became available \\citep[e.g.][]{hoflich2002,hungerford2003,maeda2006, sim2008,kromer2010}. \\citet{hoflich2002} discussed effects of the multi-dimensionality in the explosion (for a few delayed detonation models) on the line profiles of the 812 and 847 keV lines. \\citet{sim2008} examined flux evolutions in different energy bands based on kinematic models. \\citet{kromer2010} presented a prediction of the high energy signal based on the double-detonation sub-Chandrasekhar models. In this paper, we report the first study on expected high energy emission signatures and their flux evolutions from a series of two-dimensional delayed-detonation models.  We use these models to discuss possible constraints on the explosion mechanism of SNe~Ia through the $\\gamma$-ray non-detection from SN 2011fe by the currently operating instrument SPI on board {\\em INTEGRAL} \\citep[a narrow line sensitivity of $3.1 \\times 10^{-5}$ photons cm$^{-1}$ s$^{-1}$ at 1 MeV for $10^{6}$ seconds exposure:][]{roques2003}.\\footnote{We adopt the sensitivity of SPI from the latest SPI observer's manual at http://www.rssd.esa.int/ .} We then examine the detectability of the high energy emission from SNe Ia by near future observatories.  Those include {\\em NuStar} \\citep{koglin2005}, HXI \\citep[Hard X-ray Imager: ][]{kokubun2010} and SGD \\citep[Soft Gamma-ray Detector: ][]{tajima2010} on board {\\em Astro-H} \\citep{takahashi2010}, and {\\em GRIPS} \\citep{greiner2011}.  {\\em NuStar} and HXI are designed to reach to a sensitivity of a few $10^{-8}$ cm$^{-2}$ s$^{-1}$ keV$^{-1}$ (for $10^6$ seconds) in the hard X-ray range.  The SGD's designed sensitivity is (5 -- 10) $\\times 10^{-8}$ cm$^{-2}$ s$^{-1}$ keV$^{-1}$ (for $10^6$ seconds) in the soft $\\gamma$-ray range below $\\sim$600 keV.  {\\em GRIPS} is one example for proposed next generation telescopes, designed to be by a factor of $\\sim$15 better in sensitivity than {\\em INTEGRAL} in the MeV range. ",
        "conclusions": "In this paper, we have reported properties of high energy emissions from the radioactive decay chain $^{56}$Ni $\\to$ $^{56}$Co $\\to$ $^{56}$Fe in SNe Ia. A series of two-dimensional delayed-detonation models have been investigated.  We estimate, for the narrow-line sensitivity of $3 \\times 10^{5}$ photons cm$^{-1}$ s$^{-1}$, that the 847 keV line from the decay of $^{56}$Co is detectable by SPI/{\\em INTEGRAL} for the most-nearby SNe within $\\sim$6 Mpc, at 60 days after the explosion and thereafter. This, however, is likely optimistic in view of the non-detection of the signal from SN 2011fe by SPI (\\S 3.3).  The flux of the 847 keV line is sensitive to $M$($^{56}$Ni) but not to other model details, such as the progenitor mass, the flame propagation modes, 1D or multi-D, or the viewing direction.  Thus, the upper limit of the SPI observation is directly translated to a constraint on the mass of $^{56}$Ni, as $M$($^{56}$Ni) $\\lsim 1.0$ \\msun. This is not as strong as the constraint placed by optical emission analysis \\citep{nugent2011b,roepke2012}, but is totally independent from and more direct than the optical emission analysis. This shows a potential to place a strong constraint on the nature of the explosion through the high energy emission.  In the earlier phase the most constraining signal is the 158 keV line. The behavior here is sensitive to different models (e.g., the thermonuclear flame modes, initial conditions and viewing angles within the delayed-detonation scenario) -- the feature essentially traces how much material is present atop of the $^{56}$Ni-rich region.  SN 2011fe was observed by {\\em INTEGRAL} with an exposure of $\\sim$$10^6$ seconds starting at $\\sim$5 days after the discovery \\citep{isern2011a}, but unfortunately the reported upper limit is not deep enough to reject any models presented in this paper (including the pure-detonation model). For more detailed and quantitative analysis, variations in the line width predicted for different models will need to be taken into account.  While most previous studies focused on the detectability of the radioactive signals in the MeV range, we suggest that detecting soft $\\gamma$-rays and hard X-rays is more promising with the new, near-future observatories ({\\em Astro-H} and {\\em NuStar}).  We have found that the 158 keV line is detectable up to $\\sim$25 Mpc with SGD on board {\\em ASTRO-H}, and the hard X-ray continuum up to $\\sim$15 Mpc with HXI on board {\\em ASTRO-H} and {\\em NuStar}.  These near-future observatories, which we predict are able to detect the high energy emission almost annually, are expected to provide practically applicable diagnostics on the explosion mechanism. \\begin{itemize} \\item The hard X-ray continuum provides a measure of the composition near the surface. \\item The 158 keV line flux provides a measure of how much material is present above the $^{56}$Ni-rich region. \\item Accordingly, some line-to-continuum ratios, shown to be accessible by these new instruments, will provide a strong constraint on the explosion mechanism (or the viewing direction in the 2D delayed-detonation models). \\end{itemize}  Compared to other model variants (e.g., 1D pure deflagration model, 1D pure detonation model), the 2D delayed detonation model tends to predict a lower (angle-averaged) flux in these energy ranges. Thus our estimate on the detectability of these features with the future observatories may well be a conservative estimate. The hard X-ray continuum and the 158 keV line data alone could be used to obtain a rough constraint on $M$($^{56}$Ni), but this is contaminated by the factors arising from different explosion models as described above.  According to the systematic study of the 2D delayed-detonation models, we have the following solid prediction for this model sequence: For a large statistical sample, a possible scatter in the peak flux may limit the degree of the asymmetry in the explosions. We predict larger asymmetry, thus a larger scatter in the peak flux, for fainter SNe Ia according to the delayed-detonation model.  Future new generation $\\gamma$-ray observatories like {\\em GRIPS} are expected to deepen the observable horizon in the MeV range up to $\\sim$20 Mpc, or even to 35 Mpc for an exposure as long as $3 \\times 10^{6}$ seconds. The 812 keV line can be used in a way similar to the emission features in the softer band as described above. The peak 847 keV line flux alone is a very good tracer for $M$($^{56}$Ni); for the model sequences we explore in the paper (the 2D delayed-detonation models, 1D models including various flame propagation modes, those based on Chandrasekhar and sub-Chandrasekhar progenitors), it is insensitive to the model variants and the viewing direction, thus being a direct probe of explosive nucleosynthesis in SNe Ia."
    },
    "1208/1208.5488_arXiv.txt": {
        "abstract": "Recent evidence for one or more gamma-ray lines at $\\sim130$\\,GeV in the Fermi-LAT data from the Galactic Center has been interpreted as a hint for dark matter annihilation to $Z\\gamma$ or $H\\gamma$ with an annihilation cross section of $\\langle \\sigma v\\rangle$ $\\sim10^{-27}{\\rm cm^3\\,s^{-1}}$. We test this hypothesis by comparing synchrotron fluxes due to the electrons and positrons from decay of the $Z$ or the $H$ bosons {\\it only} against radio data from the same region in the Galactic Center. We find that the radio data from single-dish telescopes marginally constrain this interpretation of the claimed gamma lines for a contracted NFW profile. Already-operational radio telescopes, such as LWA, VLA-Low and LOFAR, and future radio telescopes like SKA, are sensitive to annihilation cross sections of the order of $10^{-28}{\\rm cm^3\\,s^{-1}}$, and can confirm or rule out this scenario very soon. We discuss the dependencies on the dark matter profile, magnetic fields, and background radiation density profiles, and show that the constraints are relatively robust for any reasonable assumptions. Independent of the above said recent developments, we emphasize that our radio constraints apply to all models where dark matter annihilates to $Z\\gamma$ or $H\\gamma$. ",
        "introduction": "\\label{sec:introduction} The particle identity of dark matter (DM) is one of the outstanding puzzles in contemporary physics. In order to fully understand the particle properties of dark matter, a number of complementary approaches to dark matter searches have been adopted. Indirect detection of dark matter is a promising technique, in which the products of dark matter annihilation are searched for, and gives us information about the DM abundance and annihilation rate at various astrophysical sites~\\cite{Jungman:1995df, Bertone:2004pz, Bergstrom:2012np,Feng:2010gw,Gunn:1978gr,Zeldovich:1980st}.  Gamma-ray lines from DM self annihilation are believed to be a smoking-gun signature, and have been investigated in considerable detail~\\cite{Bergstrom:1988fp, Flores:1989ru, Bergstrom:1997fh, Bern:1997ng, Ullio:1997ke, Bertone:2009cb}. Despite the relative freedom in DM model-building, if DM self-annihilation is to two-body Standard Model final states, then gamma-ray line(s) can be produced only via the following three channels{\\blue :} (i) $\\chi\\chi \\rightarrow \\gamma \\gamma$, (ii) $\\chi\\chi \\rightarrow Z \\gamma$, and (iii) $\\chi\\chi \\rightarrow H \\gamma$, where $\\chi$ denotes the DM, and $Z$ and $H$ denote the $Z$ and Higgs boson, respectively. We take the mass of the Higgs boson to be 125\\,GeV~\\cite{:2012gu, :2012gk}, and allow a heavy DM to annihilate to it.  Recently, evidence for a gamma-ray line from the Galactic Center (GC) has been uncovered in the Fermi-LAT data at $\\sim130$ GeV~\\cite{Weniger:2012tx, Bringmann:2012vr} and this has given rise to renewed interest in considering the line signal in more detail~\\cite{Tempel:2012ey,Boyarsky:2012ca, Rajaraman:2012db, GeringerSameth:2012sr, Su:2012ft, Buchmuller:2012rc, Cohen:2012me, Cholis:2012fb, Su:2012zg, Bergstrom:2012vd, Hektor:2012kc,Yang:2012ha,Huang:2012yf,Feng:2012gs,Li:2012qg,Whiteson:2012hr}.  This statistically significant signal has been tentatively interpreted as arising from DM annihilation. Generally speaking, the signal requires a DM self annihilation cross section of $\\sigv\\sim10^{-27}\\,\\rm{cm^3\\,s^{-1}}$ and the Galactic DM halo described by a standard NFW, Einasto, or a contracted NFW profile. Subsequently, a variety of particle physics models have been proposed to explain the signal~\\cite{Ibarra:2012dw, Dudas:2012pb, Cline:2012nw, Choi:2012ap, Kyae:2012vi, Lee:2012bq, Acharya:2012dz, Buckley:2012ws, Chu:2012qy, Das:2012ys, Kang:2012bq, Weiner:2012cb, Heo:2012dk,Tulin:2012uq,Li:2012jf,Park:2012xq,Frandsen:2012db,Oda:2012fy,Cline:2012bz,Bai2012,Shakya:2012fj,Bergstrom:2012bd,Fan:2012gr}. It is also found that the line is off-center from the GC by approximately $1.5^\\circ$~\\cite{Su:2012ft, Yang:2012ha}, which requires the center of the DM halo to be displaced from the baryonic center. This degree of displacement appears reasonable as shown by recent numerical simulations~\\cite{Kuhlen:2012qw}.  On the other hand, there are arguments against the DM origin of the gamma-ray line. There are hints that the line is also present in the photons collected from the cascades in the Earth's atmosphere, which is a ``pure background\" region~\\cite{Su:2012ft}, although this has been claimed to be due to statistical fluctuations~\\cite{Finkbeiner:2012ez,Hektor:2012ev}. There have been claims of the presence of gamma-ray lines at the same energy, spatially correlated with some Fermi-LAT unassociated sources~\\cite{Su:2012zg}. However, there are also counterclaims that most of these unassociated sources are consistent with being standard astrophysical objects such as active galactic nuclei or statistical fluctuations~\\cite{Hooper:2012qc,Mirabal:2012za,Hektor:2012jc}. Furthermore, it remains possible that the GC line signal is also of an astrophysical origin~\\cite{Boyarsky:2012ca,Profumo:2012tr, Aharonian:2012cs}.  The Fermi-LAT collaboration, in their search for $\\gamma$-ray lines in the 2 year data set~\\cite{Ackermann:2012qk} did not find a signal as the analysis employed a  different search strategy, an older data set and background rejection software, and a larger search region, making it difficult to compare directly with the above claims. However in their most recent search for gamma-ray lines with the 4 year data~\\cite{Albert:Fermi2012}, the Fermi-LAT collaboration has acknowledged the presence of a feature at the GC at 135 GeV (this shift in the energy is due to recalibration but we will assume that the line is at 130 GeV throughout this work). The collaboration also finds a feature at the Earth limb at the same energy~\\cite{Charles:Fermi2012}. The collaboration states that it does not have a consistent interpretation of the Galactic Center feature and that it needs more data to resolve the issue~\\cite{Bloom:Fermi2012}. Given the arguments in favor of and against the DM origin of this signal, this remains a topic of active research.  If the DM annihilates to two-body Standard Model final states, as in (i)-(iii), then we can predict some particle physics model-independent consequences. For a dominant annihilation (i), i.e., to two photons, there are no further interactions of the photons at an appreciable level, with all higher-level amplitudes suppressed by at least $\\alpha~\\approx~1/137$. However, if the annihilation proceeds as~in~(ii) or (iii), i.e., to a photon and a heavy Standard Model boson, the heavy boson decays to other charged particles which can have observable consequences.  The decay of the $Z$ and the $H$ boson produces electrons, protons, neutrons, neutrinos, their antiparticles, as well as photons as final states. The almost featureless spectra of these secondary particles poses considerable difficulty in their search above the astrophysical backgrounds. Searches for antimatter benefit from lower cosmic ray backgrounds, therefore one can search for antiprotons and positrons from the $Z$ and the $H$ boson. A search for antiprotons from these decays constrains several particle physics models which can give rise to a gamma-ray line~\\cite{Buchmuller:2012rc}, whereas the preexisting unaccounted excess in positrons~\\cite{Adriani:2008zr, FermiLAT:2011ab} makes a positron search ambiguous. Neutrinos could, in principle, be used to distinguish between all three final states, but achieved or projected sensitivities in the range \\mbox{$\\sigv\\sim(10^{-22}\\,{\\rm cm^3\\,s^{-1}} - 10^{-23}\\,{\\rm cm^3\\,s^{-1}}$)~\\cite{Abbasi:2011eq,Beacom:2006tt,Yuksel:2007ac,Dasgupta:2012bd}} will not be able to probe the claimed signal. Secondary photons that are produced in the decay of the $Z$ or the $H$ boson, or in other DM annihilation channels, also constrain these scenarios~\\mbox{\\cite{Buchmuller:2012rc, Cohen:2012me, Cholis:2012fb}}, and there are ongoing efforts to confirm this 130\\,GeV line with future detectors~\\cite{Bergstrom:2012vd,Li:2012qg}.  In this paper, we ask the questions -- If the 130\\,GeV signal is indeed from DM annihilation to $Z \\gamma$ or $H \\gamma$, what other consequences are guaranteed? Can we use these consequences to test this signal? Synchrotron radiation from products of DM annihilation has been argued to provide strong constraints for many DM annihilation channels and scenarios~\\cite{Aloisio:2004hy, Borriello:2008gy, Bertone:2008xr, Ishiwata:2008qy, Fornengo:2011iq, Crocker:2010gy, Gondolo:2000pn, Bergstrom:2008ag, Hooper:2007kb, Regis:2008ij, Bertone:2001jv, Hooper:2010im, Hooper:2012jc, Fornengo:2011cn, Zhang:2009pr, Zhang:2008rs,Mambrini:2012ue,Fornengo:2011xk}. Thus, following these promising leads, we explore our question by calculating the synchrotron radiation in the GC due to the electrons and positrons from $Z$ or $H$ decays {\\it only}, and comparing it to existing data from radio telescopes.  We first take a very conservative approach, where we compare the DM-induced synchrotron fluxes to the total measured radio flux at 330\\,MHz in a relatively large region around the GC, and determine that DM annihilation cross sections to these channels cannot be more than $\\sigv\\sim10^{-25}{\\rm cm^3\\,s^{-1}}$. However, this approach is overly conservative, as the synchrotron fluxes in the GC are modeled accurately with known astrophysics. We argue that the flux due to dark matter must not exceed the uncertainties on the modeled radio fluxes, which provides us with a constraint $\\sigv \\sim10^{-26}{\\rm cm^3\\,s^{-1}}$. Constraints obtained by comparing fluxes predicted in smaller regions of interest and upper limits at 408\\,MHz imply $\\sigv\\sim10^{-27}{\\rm cm^3\\,s^{-1}}$, and are already in mild tension with the 130\\,GeV line. We forecast that the sensitivity can be improved to $\\sim10^{-28}{\\rm cm^3\\,s^{-1}}$ with a few hours of observation of the GC at 80 MHz with LWA, the Long Wavelength Array, and at 200\\,MHz with LOFAR, the LOw-Frequency ARray for radio astronomy, allowing us to constrain interpretations of the 130\\,GeV line signal in the Fermi-LAT data in terms of DM annihilation to $Z \\gamma$ or $H \\gamma$. Although this is the main motivation for our present work, the radio constraints we derive are valid regardless of whether this claimed 130\\,GeV line signal survives further scrutiny or not. These constraints will continue to apply to any future interpretations of gamma-ray lines at the GC in terms of DM annihilation.  Note that, these sensitivities readily probe the cross section that explains the tentative 130\\,GeV line signal. More generally, we expect these sensitivities to be able to probe many of the models, not necessarily supersymmetric, that could explain this signal. We also emphasize that since we are looking for the synchrotron radiation from the electrons and positrons produced in the decays of the $Z$ or the $H$ boson {\\it only}, our constraints are independent of the underlying DM particle physics model. In these two ways, our work is complementary to Ref.\\,\\cite{Cohen:2012me}. The results here are of course affected by astrophysical uncertainties, e.g., dark matter density profile, magnetic fields, interstellar radiation energy density, and proton density in the Galaxy, and by taking a range of different values for them we try to understand their impact.  The rest of the paper is organized as follows. In Sec.\\,\\ref{sec:DataTheory} we discuss the radio data that we use for obtaining our constraints, and the theoretical framework for calculating the flux densities from synchrotron radiation by DM annihilation products. In Sec.\\,\\ref{sec:Inputs} we furnish and justify the astrophysical inputs, i.e., DM density, magnetic fields, and radiation density in the Galaxy, that we use for our calculations. In Sec.\\,\\ref{sec:results} we show the predicted flux densities for benchmark DM annihilation cross sections, and provide constraints on the DM annihilation cross section as a function of DM mass for the channels \\mbox{$\\chi\\chi\\rightarrow Z\\gamma$} and $\\chi\\chi\\rightarrow H\\gamma$, and conclude in Sec.\\,\\ref{sec:outlook}. ",
        "conclusions": "\\label{sec:outlook} In this paper we have shown that existing radio data around the Galactic Center at 408\\,MHz marginally constrains the interpretation of the 130\\,GeV line in Fermi-LAT data in terms of DM self annihilation to $Z\\gamma$ or $H\\gamma$ with a cross section $\\sim$ $\\,10^{-27}\\,{\\rm cm^3\\,s^{-1}}$ for a contracted NFW profile. For other frequencies or other DM density profiles the constraint is up to an order of magnitude weaker within the parameter ranges chosen by us. Future measurements made around the GC by LWA in the 80\\,MHz band can push the sensitivity to DM annihilating to gamma-ray lines down to $\\sigv\\sim10^{-28}\\,{\\rm cm^3\\,s^{-1}}$ and enable a test of the above signal. Although the background needs to be known very well to achieve our quoted limits, these possibilities are, to the best of our knowledge, some of the most competitive ways to test for the nature of the DM that could have produced the tentative 130\\,GeV line signal.  We have shown that these conclusions are fairly robust with respect to the assumptions on the magnetic field in the Galaxy, and the constraints do not weaken by more than an order of magnitude. The dependence on DM density profiles is somewhat more important, especially when the region of observation is small and closely centered on the GC. While the uncertainty in the astrophysical modeling of the GC does impact our results (see for e.g., \\cite{Everett:2007dw} for a different modeling of the GC), we must emphasize that these constraints are completely model-independent from the particle physics perspective, because we have simply taken the electrons and positrons from the known decays of the $Z$ or $H$ produced in the DM annihilation to $Z\\gamma$ or $H\\gamma$, respectively. A similar study on dark matter annihilation contribution to the galactic radio background~\\cite{Kogut:2009xv} and diffuse extragalactic radio background~\\cite{Fixsen:2004hp,Seiffert:2009xs} can also performed to cross-check potential dark matter signals from the Galactic Center~\\cite{Hooper:2012jc,Fornengo:2011cn,Fornengo:2011xk}.  We hope that these results will encourage radio astronomers, especially those at LWA, VLA-Low, LOFAR, and SKA, to observe the GC, model the astrophysical synchrotron backgrounds, and determine if there is any excess flux. Irrespective of whether the tentative 130\\,GeV gamma-ray line signal at Fermi-LAT is due to DM annihilation or not, this promises to deliver some of the strongest constraints on DM annihilation."
    },
    "1208/1208.1112_arXiv.txt": {
        "abstract": "The correlation between infrared-to-ultraviolet luminosity ratio and ultraviolet color (or ultraviolet spectral slope), i.e., the IRX-UV (or IRX-$\\beta$) relation, found in studies of starburst galaxies is a prevalent recipe for correcting extragalactic dust attenuation. Considerable dispersion in this relation discovered for normal galaxies, however, complicates its usability. In order to investigate the cause of the dispersion and to have a better understanding of the nature of the IRX-UV relation, in this paper, we select five nearby spiral galaxies, and perform spatially resolved studies on each of the galaxies, with a combination of ultraviolet and infrared imaging data. We measure all positions within each galaxy and divide the extracted regions into young and evolved stellar populations. By means of this approach, we attempt to discover separate effects of dust attenuation and stellar population age on the IRX-UV relation for individual galaxies. In this work, in addition to dust attenuation, stellar population age is interpreted to be another parameter in the IRX-UV function, and the diversity of star formation histories is suggested to disperse the age effects. At the same time, strong evidence shows the necessity of more parameters in the interpretation of observational data, such as variations in attenuation/extinction law. Fractional contributions of different components to the integrated luminosities of the galaxies suggest that the integrated measurements of these galaxies, which comprise different populations would weaken the effect of the age parameter on IRX-UV diagrams. The dependance of the IRX-UV relation on luminosity and radial distance in galaxies presents weak trends, which offers an implication of selective effects. The two-dimensional maps of the UV color and the infrared-to-ultraviolet ratio are displayed and show a disparity in the spatial distributions between the two parameters in galaxies, which offers a spatial interpretation of the scatter in the IRX-UV relation. ",
        "introduction": "Obscuration of starlight by interstellar dust grains is a serious obstacle which hampers our ability to directly derive stellar population properties of galaxies from observed radiative information. Correction for dust attenuation is crucial for determining galactic properties such as star formation history (SFH) in current extragalactic astronomy. Calibration of observational indicators to intrinsic parameters of galaxies, for instance from observed UV-optical flux to star formation rate (SFR), is in tight connection with the fact of proper compensation made for dust attenuation \\citep[e.g.,][]{2004A&A...425..417K, 2009ApJ...703.1672K, 2009ApJ...706..599L, 2011ApJ...741..124H}.  Dust grains absorb stellar emission with the wavelength coverage from far-ultraviolet to near-infrared, and re-emit the energy as mid- and far-infrared thermal radiation. According to the energy balance theory, the luminosity ratio of total infrared (IR) to ultraviolet \\citep[UV; so-called IR excess or IRX, initially introduced in][]{1974A&A....32..269M} has been invented and considered as a reliable indicator of dust attenuation \\citep{1992A&A...264..444B, 1995A&A...293L..65X, 1996A&A...306...61B, 2000ApJ...528..799W, 2005ApJ...619L..51B}. Throughout this paper, the IRX is defined in the form of a logarithm: $\\mathrm{IRX} \\equiv \\log(L(\\mathrm{IR})/L(\\mathrm{FUV}))$, where L(IR) is luminosity of infrared thermal radiation of dust and L(FUV) is far-ultraviolet monochromatic luminosity of starlight (detailed definitions of these terms are presented in Section \\ref{measure}). In the meanwhile, for starburst galaxies, it is believed that the ultraviolet spectral slope \\citep[$\\beta$, defined as $f_\\lambda \\propto \\lambda^{\\beta}$, where $1250 \\mathrm{Ang} \\leq \\lambda \\leq 2600 \\mathrm{Ang}$;][]{1994ApJ...429..582C} has an intrinsic value \\citep[$-2.3 \\lesssim \\beta \\lesssim -2.0$, as suggested in][]{1994ApJ...429..582C, 1995ApJS...96....9L, 1999ApJ...521...64M}, and any change in this slope results from wavelength-selective absorption of photons by dust grains. On a basis of this assumption, the slope of UV spectrum is regarded as another attenuation estimator \\citep{1994ApJ...429..582C, 1997AJ....113..162C}.  The correlation between the two parameters was found by \\citet[][so-called IRX-$\\beta$ relation]{1999ApJ...521...64M} with an investigation based on a local starburst sample, and the relation was believed to be a sequence indicating dust attenuation. Therefore, the IRX-$\\beta$ relation has great significance for attenuation correction, since it provides an access to estimate dust attenuation solely with UV waveband. This prescription has been widely used at the present time especially for high redshift galaxies, of which the rest-frame UV observations are available from ground-based optical telescopes \\citep[e.g.,][]{2004ApJ...617..746D, 2007ApJ...670..156D, 2006ApJ...638...72K, 2007ApJS..173..415M, 2009ApJ...705..936B}. Since people usually adopt UV color or UV luminosity ratio as a surrogate for the traditional UV spectral slope, hereafter we refer to the relation as the terminology \"IRX-UV\" in place of the \"IRX-$\\beta$\".  Observations on extragalactic star-forming regions and normal galaxies following the discovery of the IRX-UV relation show considerable dispersion in this relation. The analysis of HII regions in the Large Magellanic Cloud by \\citet{2002ApJ...577..150B} and the study of normal galaxies by \\citet[][hereafter denoted as K04]{2004MNRAS.349..769K} have shown that the IRX-UV relation appears to have flatter UV spectral slopes than starburst galaxies at fixed IRX and covers a wide range; the tight correlation between IRX and $\\beta$ for starburst galaxies has been found in absence, with a considerable degree of deviation and dispersion from the starburst formula instead.  Following the launch of the \\emph{Galaxy Evolution Explorer} \\citep[\\emph{GALEX};][]{2003lgal.conf...10B, 2005ApJ...619L...1M} and the \\emph{Spitzer Space Telescope} \\citep[\\emph{Spitzer};][]{2004ApJS..154....1W}, an increasing number of studies have revealed the deviation in the IRX-UV relation, from statistics of integrated measurements of galaxies as a whole \\citep[e.g.,][]{2007ApJ...655..863D, 2009ApJ...703..517D, 2007ApJS..173..185G} to spatially resolved studies of individual galaxies \\citep[e.g.,][]{2004ApJS..154..215G, 2005ApJ...633..871C, 2007ApJS..173..572T}. It is worth mentioning that, with the combination of larger and more comprehensive UV and IR all sky survey data, \\citet{2005ApJ...619L..51B} have found the similar offset and scatter from the starburst relation in the IRX-UV diagram, while a few number of IR-bright galaxies appear to have higher IRXs than starburst galaxies at given UV spectral slopes. The IRX-rising phenomenon has been confirmed by studies of (ultra-)luminous infrared galaxies, and adds an additional question to the IRX-UV issue \\citep{2002ApJ...568..651G, 2010ApJ...715..572H}. In order to improve the IRX-UV calibration to dust attenuation and to provide better insights into the nature of galactic UV and IR properties, it is necessary to discover the origin of the deviation in the IRX-UV relation.  A variety of SFH of different stellar populations have been considered as one main cause of the deviation, and in simple cases this deviation is indicated by stellar population age. In K04 work, the authors have compiled a sample of galaxies with a wide coverage of star formation activities from starburst to quiescent, and employed the spectral index $D_n(4000)$ to trace different stellar populations. Via the analysis of perpendicular distances of the galaxies on the IRX-UV diagram from the empirical best-fitting curve for starburst galaxies (the starburst empirical curve is parameterized by Equation (2) in K04 paper), they have confirmed that variations in stellar populations dominate the dispersion in the IRX-UV relation. The products of stellar population synthesis modeling in K04 work have shown that aging of stellar populations would yield redder unattenuated UV colors (i.e., flatter intrinsic UV spectral slopes), and the effect of stellar population age on the IRX-UV relation is depicted as a series of sequences in parallel with the starburst empirical relation (Figure 4 in K04 paper). Thus, in addition to dust attenuation, stellar population age has been considered as the second parameter in the IRX-UV function. Nevertheless, there has been lack of strong evidence to demonstrate such an age effect, and in many cases it is ascribed to the fact that the adopted observational tracers are inappropriate, or the age effects are weak compared with other potential effects and hence can be easily masked \\citep{2005ApJ...619L..55S, 2006ApJ...637..242C, 2007ApJS..173..392J}. Despite these disappointments, studies of radial profiles for individual galaxies have disclosed a weak trend of the age-sensitive color FUV $-$ 3.6 $\\mu$m in the IRX-UV relation, which is in encouraging support of the role stellar population age plays in the IRX-UV relation \\citep{2009ApJ...701.1965M}. Notwithstanding, the radial profiles still cannot adequately resolve the mixing of stellar populations within galaxies.  A recent work on a basis of pixel-by-pixel studies of nearby galaxies taken by \\citet{2012A&A...539A.145B} makes use of the data obtained from the \\emph{Herschel Space Observatory} \\citep{2010A&A...518L...1P} to examine a set of parameters. The results in their work have illustrated that, intrinsic UV color plays a predominant role in the IRX-UV relation; the commonly used age indicators including the spectral index $D_n(4000)$ are in a poor position to estimate accurate stellar population age and therefore fail to trace the IRX-UV relation as a function of age; in addition, the analysis in \\citet{2012A&A...539A.145B} has further predicted a non-negligible contribution from the shape of attenuation law to the deviation in the IRX-UV relation. Coincidentally, several numerical simulations have reproduced a diversity of IRX-UV trends by varying dust-star geometrical configuration and dust grain properties \\citep{2005MNRAS.360.1413B, 2007MNRAS.375..640P, 2008MNRAS.386.1157C}.  In this work, we carry out spatially resolved studies of galaxies which enable an inspection of galactic sub-structures and allow us to extract certain sources of interest within each single galaxy rather than a mixture of different stellar populations \\citep[e.g.,][]{2005ApJ...633..871C, 2006ApJ...648..987P, 2007ApJS..173..572T}, and in this way we can have better separation between different populations than integrated measurements or radial profiles of galaxies. This kind of investigation also makes it possible to compare regions of a similar stellar population in different galaxies. \\citet{2009ApJ...706..553B} have been engaged in a census of star-forming clusters in eight nearby spiral galaxies. Further to the approach taken by \\citet{2009ApJ...706..553B}, we focus on all positions in galaxies, not only young clusters, but also evolved stellar populations in galactic background areas which are supposed to have different intrinsic UV colors from young stellar populations. This strategy is a natural way to define stellar populations and makes it possible to exhibit age signatures in the IRX-UV diagrams. Up to now, most analyses are based on integrated measurements of galaxies, while in this work, we take advantage of the spatially resolved analysis to explore the determinant of the locations of integrated galaxies in the IRX-UV diagram. Also, we test the systematic dependence of the IRX-UV relation on FUV luminosity and radial distance, and present two-dimensional maps of UV color and IRX to provide a spatial insight into the IRX-UV relation.  The remainder of this paper is outlined as follows. In Section \\ref{data}, we describe the compilation of the sample, the data processing and the photometric measurements, and the modeling of spectral synthesis; in Section \\ref{result1}, we present the resulting IRX-UV diagrams for the extracted regions within the galaxies in our sample; in Section \\ref{result2}, we examine the systematic dependance of the IRX-UV relations on luminosity and radial distance; in Section \\ref{2-D map}, we provide spatial distributions of UV color and IRX; Section \\ref{disc} is concerned with relevant discussion and interpretation about the presented results; finally, we summarize the results and their implications in Section \\ref{sum}. ",
        "conclusions": "\\label{disc}  The main goal of this work is to have a better understanding of the two parameters (dust attenuation and stellar population age) in the IRX-UV function. In Section \\ref{result1}, we present the spatially resolved results of the IRX-UV diagrams. We can see the age effects appearing as the displacements between the locations of the UV clusters and the local background regions inside galaxies in these diagrams, and the deviations of the age parameter in the interpretations of the IRX-UV distributions. The two-parameter scenario adopted for characterizing the observational data in Section \\ref{result1} assumes simple stellar populations with an instantaneous burst. This assumption is an approximation of stellar populations contained in galactic subregions. As a matter of fact, although subregions inside galaxies are simpler in SFH than integrated galaxies, stellar populations still tend to be born beyond the instantaneous burst on a scale of several hundreds parsecs, which has the potential to oversimplify the stellar populations and introduce inaccuracies in the resulting descriptions of age for the measured regions. In this section, we will propose a series of composite stellar populations for the modeled scenarios, and discuss potential influences of the complexities in SFH on the IRX-UV properties.  Spectra of the composite stellar populations are constructed via the GALAXEV library of evolutionary population synthesis \\citep{2003MNRAS.344.1000B}, on an assumption of exponentially decreasing SFR ($\\sim e^{-t/\\tau_\\mathrm{SF}}$, where $t$ is stellar population age and $\\tau_\\mathrm{SF}$ is constant of star formation timescale), with the \\citet{2003PASP..115..763C} IMF \\footnote{The GALAXEV spectral library adopts the Chabrier IMF in stellar population synthesis, while the STARBURST99 spectral library offers the Kroupa IMF. Both IMFs are essentially identical for stellar mass $\\geq 1 M_\\odot$. The difference between the two IMFs exists in descriptions of low-mass stars and brown dwarfs, and has no effective influence on UV stellar emission studied in this work.} and solar metallicity ($Z=0.02$). Figure \\ref{color_vs_age} shows intrinsic $\\mathrm{FUV}-\\mathrm{NUV}$ as a function of stellar population age in various SFH scenarios: simple stellar populations with an instantaneous burst, composite stellar populations with six types of exponential decreases in SFR ($\\tau_\\mathrm{SF}$ = 0.01, 0.1, 0.5, 1, 2, and 8 Gyr), and constant SFR of 1 $M_\\odot~yr^{-1}$. We produce two scenarios of an instantaneous burst in this figure from STARBURST99 and GALAXEV respectively for the purpose of comparison between the two libraries. The exponentially decreasing and constant SFRs are derived from the GALAXEV library. As clearly illustrated in this figure, stellar populations born with an instantaneous burst present the most tremendous reddening of UV color with increasing age, in particular at age $>$200 Myr: $\\mathrm{FUV}-\\mathrm{NUV}$ extends from 0.3 to $\\sim$3.0 mag within an age range of 200$-$700 Myr. The reddening evolution of UV color tends to fade in the scenarios of composite stellar populations, and it is apparent to see at longer star formation time that $\\mathrm{FUV}-\\mathrm{NUV}$ appears to be more insensitive to age. The exponentially decreasing SFR with $\\tau_\\mathrm{SF}$ = 0.01 Gyr has a color evolution comparable to the instantaneous burst; at $\\tau_\\mathrm{SF}$ = 0.1 Gyr, the reddening of UV color becomes intensive after 500 Myr age; for longer star formation time, $\\mathrm{FUV}-\\mathrm{NUV}$ retains a constant value $\\sim$0.0 mag until stellar populations evolve to $\\sim$2 Gyr at $\\tau_\\mathrm{SF}$ = 0.5 Gyr, $\\sim$5 Gyr at $\\tau_\\mathrm{SF}$ = 1 Gyr, and $\\sim$10 Gyr at $\\tau_\\mathrm{SF}$ = 2 Gyr; stellar populations with $\\tau_\\mathrm{SF}$ = 8 Gyr and the constant SFR produce identical evolutionary tracks of UV color, where $\\mathrm{FUV}-\\mathrm{NUV}$ $\\sim$ 0.0 mag during the lifetime up to 20 Gyr. Complex SFHs introduce the degeneracy of stellar population age and star formation timescale in $\\mathrm{FUV}-\\mathrm{NUV}$: larger $\\tau_\\mathrm{SF}$ means more extended SFH and results in bluer color than smaller $\\tau_\\mathrm{SF}$ at constant age.  It is worth noting in Figure \\ref{color_vs_age} that, there is a reversing point on each of the evolutionary tracks produced with the GALAXEV library except the SFHs of $\\tau_\\mathrm{SF}$ = 8 Gyr and the constant SFR, which suggests an upper limit for UV color at a certain age, and after this stage the color is supposed to experience a kind of blueing evolution. This feature is different from the STARBURST99 product where $\\mathrm{FUV}-\\mathrm{NUV}$ is on the monotonic increase during the whole lifetime of stellar populations, by comparison between the two models for the same instantaneous burst in this figure. It is also displayed that the reversing points on the GALAXEV tracks emerge in later periods and the upper limits appear at lower values if stellar populations evolve with larger $\\tau_\\mathrm{SF}$. As shown in this figure, the SFHs of an instantaneous burst and $\\tau_\\mathrm{SF}<0.1$ Gyr yield the reddest colors over 3.6 mag during 1$-$2 Gyr; for $\\tau_\\mathrm{SF}$ = 2 Gyr, the maximum in $\\mathrm{FUV}-\\mathrm{NUV}$ drops to $\\sim$0.3 mag, and the time falls to $\\sim$20 Gyr; and between them, the reversing point for longer star formation time appears to locate on the track for lower $\\tau_\\mathrm{SF}$, and after the reversing point the both SFHs tend to share a common evolutionary track. The presence of the blueing evolution is ascribed to the additional account of FUV emission from late-type post-asymptotic-giant-branch (post-AGB) stars in the GALAXEV library \\citep{2003MNRAS.344.1000B}. Figure \\ref{comp_model} shows a comparison between UV spectral energy distributions (SEDs) derived respectively from the STARBURST99 and GALAXEV libraries for an instantaneous burst at two speci\ufb01c ages as an example to illustrate the difference in spectral shape. The contribution of the post-AGB stars manifests as the constant spectrum with an even rising trend shortwards from 2000 $\\mathrm{Ang}$ wavelength derived from the GALAXEV library for 8 Gyr age; whereas the STARBURST99 spectrum for this age presents a steep decline in the same wavelength range.  In Figure \\ref{IRXUV_UV_M81_SFH}, we superimpose grids reproduced with composite stellar populations on the IRX-UV diagrams for NGC~3031 as plotted in Figure \\ref{IRXUV_UV_M81}, and four types of exponential decreases in SFR ($\\tau_\\mathrm{SF}$ = 0.1, 0.5, 1, and 8 Gyr) are shown respectively in each panel. The modeled curves with constant ages are displayed in parallel with each other in the figure, and with increasing ages the equivalent amounts of dust attenuation tend to stand at moderately higher levels of IRX. Compared with the SFH with an instantaneous burst, these curves for composite stellar populations are confined in narrower $\\mathrm{FUV}-\\mathrm{NUV}$ spaces, and the IRX-UV functions with different ages tend to approach each other with increasing $\\tau_\\mathrm{SF}$. For instance, in the IRX-UV planes, an age interval between 2 and 500 Myr for $\\tau_\\mathrm{SF} = 0.1$ Gyr introduces a difference of $\\sim$0.7 mag in $\\mathrm{FUV}-\\mathrm{NUV}$, while in the same age scale but for $\\tau_\\mathrm{SF} \\geq 0.5$ Gyr, the color differences diminish to less than 0.4 mag. In addition to the five certain ages addressed in Section \\ref{result1} (Figures \\ref{IRXUV_UV_M81}$-$\\ref{IRXUV_UV_total}), we further sample three older ages in Figure \\ref{IRXUV_UV_M81_SFH}: 800 Myr, 3 Gyr, and 8 Gyr. In the panel of $\\tau_\\mathrm{SF} = 0.1$ Gyr, the curve with the age of 800 Myr lies in a range where $\\mathrm{FUV}-\\mathrm{NUV} > 1.5$ mag;\\footnote{In this panel, the curve with the age of 8 Gyr suffers from the color reversing as we have presented in the above paragraph, and due to this factor this curve appears on a bluer location than that of 800 Myr.} while for $\\tau_\\mathrm{SF} \\geq 0.5$ Gyr, the same regimes in the IRX-UV diagrams belong to stellar populations evolving over 8 Gyr; when $\\tau_\\mathrm{SF}$ increases to 8 Gyr, ages ranging from 100 Myr to 8 Gyr are not distinguishable in the IRX-UV planes, and stellar populations in this timescale have intrinsic colors $\\mathrm{FUV}-\\mathrm{NUV} \\sim 0.0$ mag.  The model with composite stellar populations manifests substantial disparities in description of the same observational data with different SFH scenarios. In our work, the UV clusters inside the galaxies are measured with the subtraction of underlying diffuse emission (see Section \\ref{measure}), which enables the extraction stellar populations with an instantaneous burst or short star formation timescale, and $\\tau_\\mathrm{SF} > 0.1$ Gyr is not considered to be appropriate for the UV clusters; by contrast, the local background populations are supposed to experience longer-term star formation in SFH. In this situation, we expect to adopt diverse SFH scenarios to characterize the two populations.  The IRX-UV diagram for NGC~3031 shows a clear separation between the UV clusters, the disk background regions, and the bulge background regions (Figure \\ref{IRXUV_UV_M81}). As we have noted in Section \\ref{3031}, a considerable number of the disk background regions are categorized in the same age range for the UV clusters in the scenario of simple stellar populations, which is one of the discrepancies between the observational data and the model scenario in this paper. To address this problem, the model of composite stellar populations offers a suggestion from the viewpoint of SFH, where the local background regions are characterized as older stellar populations with increasing star formation timescales. In Figure \\ref{IRXUV_UV_M81_SFH}, for instance, the scenario with $\\tau_\\mathrm{SF} = 0.1$ Gyr offers a description of age $<$1 Gyr for the local background regions, while in the scenario with $\\tau_\\mathrm{SF} = 0.5$ Gyr most of these regions correspond to stellar population age extending to 8 Gyr. Due to the variations in SFH, it is very likely for stellar populations with a large disparity in age to lie close to each other or even overlap in the IRX-UV diagram, and in this case, the appearance of the age parameter tends to be of less prominence in the IRX-UV function.  \\begin{figure} \\centering \\vspace*{-10mm} \\includegraphics[width=\\columnwidth]{color_vs_age} \\vspace*{-55mm} \\caption{$\\mathrm{FUV}-\\mathrm{NUV}$ color as a function of stellar population age for different SFHs: simple stellar populations with instantaneous burst, composite stellar populations with exponential decreases in SFR, and constant SFR. The instantaneous burst functions are derived from two models of stellar population synthesis: STARBURST99 (black solid line) and GALAXEV (gray dotted line), for comparison. The exponentially decreasing SFRs with time constants of $\\tau_\\mathrm{SF}$ = 0.01 (red solid line), 0.1 (orange solid line), 0.5 (yellow solid line), 1 (green solid line), 2 (cyan solid line), and 8 Gyr (blue solid line), and the constant SFR of 1 $M_{\\odot}~yr^{-1}$ (purple solid line) are derived from GALAXEV synthesis model. The simple and composite stellar populations are modeled with solar metallicity ($Z=0.02$). The middle panel is plotted to show the reversing points on the modeled tracks with $\\tau_\\mathrm{SF} \\leq 0.1$ Gyr.}\\label{color_vs_age} \\end{figure}  \\begin{figure} \\centering \\vspace*{-10mm} \\includegraphics[width=\\columnwidth]{comp_model} \\vspace*{-55mm} \\caption{Top: Modeled spectra in UV wavelength range of stellar populations with instantaneous burst, $10^5 M_{\\odot}$, and solar metallicity ($Z=0.02$). The spectra are color-coded by stellar population age and relevant models: 10 Myr from GALAXEV model (orange), 10 Myr from STARBURST99 model (blue), 8 Gyr from GALAXEV model (red), and 8 Gyr from STARBURST99 model (cyan). Bottom: Filter transmission curves of two \\emph{GALEX} bandpasses: FUV (magenta) and NUV (green).}\\label{comp_model} \\end{figure}  \\begin{figure*} \\centering \\vspace*{-10mm} \\includegraphics[width=1.8\\columnwidth]{IRXUV_tot_UV1_M81_SFH} \\vspace*{-90mm} \\caption{IRX vs. $\\mathrm{FUV}-\\mathrm{NUV}$ for NGC~3031, and superimposed is the grid modeled with exponentially decreasing SFRs with time constants of $\\tau_\\mathrm{SF}$ = 0.1 (top left panel), 0.5 (top right panel), 1 (bottom left panel), and 8 Gyr (bottom right panel). Symbols are the same with those in Figure \\ref{IRXUV_UV_M81}. Black dashed lines describe the model curves sampled with five ages: 2, 8, 100, 300, and 500 Myr, from left to right on the horizontal axis, corresponding to those in Figures \\ref{IRXUV_UV_M81}$-$\\ref{IRXUV_UV_NGC7331}. Solid lines indicate older ages and are color-coded by 800 Myr (cyan), 3 Gyr (yellow), and 8 Gyr (orange). Dotted lines connect the points of five constant amounts of dust attenuation ($A_\\mathrm{V}$ = 0.01, 0.1, 0.5, 1.0, and 2.0) on the model curves of different ages, corresponding to those in Figures \\ref{IRXUV_UV_M81}$-$\\ref{IRXUV_UV_NGC7331}. Grey dot-dashed line is K04 starburst curve. Error bars showing the photometric uncertainties are plotted as well.}\\label{IRXUV_UV_M81_SFH} \\end{figure*}  The separation between the UV clusters and the local background regions in the IRX-UV diagram is obviously displayed for NGC~4536 and NGC~7331 (aside from the ring area), visible but not very apparent to distinguish for NGC~6946, and completely absent for NGC~5194 (Figures \\ref{IRXUV_UV_M51}$-$\\ref{IRXUV_UV_NGC7331}). When the diversity of SFH is applied to these galaxies, the local background regions are expected to have longer star formation timescales than the UV clusters, and the factual difference in stellar population age between the two components inside each galaxy tends to exceed the interpretation of the offset in the IRX-UV diagram on assumption of one identical SFH. Combining the data points in Figures \\ref{IRXUV_UV_M51}$-$\\ref{IRXUV_UV_NGC7331} and the modeled curves in Figure \\ref{IRXUV_UV_M81_SFH}, we can see that, most of the local background regions are characterized to cover an age range of 1$-$8 Gyr with $\\tau_\\mathrm{SF} = 0.5$ Gyr, and longer star formation timescales assign the same age range for bluer color spaces. In such cases, the UV clusters and the local background regions are able to possess the adjacent or even common regimes in $\\mathrm{FUV}-\\mathrm{NUV}$ at similar levels of IRX, in spite of possible disparities in stellar population age.  However, the variations in SFH still fail to provide any better solution to the other discrepancies in characterizing the IRX-UV relation for the galaxies with the scenario of simple stellar populations presented in Section \\ref{result1}. It is evident that, a few number of UV clusters inside NGC~3031 and several local background regions inside NGC~5194 which have bluer UV colors at fixed IRX than the prediction of the scenario with an instantaneous burst remain outside the coverage of the grid modeled with any other SFH, and a half of UV clusters inside NGC~5194 and ring clusters inside NGC~7331 which are estimated as several hundreds of megayears assuming an instantaneous burst possess even older regimes in age with increasing star formation timescales.  Figure \\ref{IRXUV_UV_total} shows significant dispersion in the composite IRX-UV relation for a total number of the UV clusters, and the scenario with an instantaneous burst fails to provide an adequate interpretation of the data distribution. The variations in SFH are shown in Figure \\ref{IRXUV_UV_M81_SFH} to have an effective impact on the IRX-UV properties for stellar populations with age $\\geq$100 Myr. For instance, between the SFHs with an instantaneous burst and the $\\tau_\\mathrm{SF} = 0.1$ Gyr star formation, there is a scatter of $\\sim$0.3 mag in $\\mathrm{FUV}-\\mathrm{NUV}$ at fixed IRX for the age of 100 Myr, and a scatter of $\\sim$0.6 mag in $\\mathrm{FUV}-\\mathrm{NUV}$ for 300 Myr. Older ages and longer star formation timescales are not taken into consideration for the UV clusters. In this situation, although UV clusters are believed to contain much simpler stellar populations than galactic background, the limited variations in SFH are still likely to introduce a moderate level of dispersion in Figure \\ref{IRXUV_UV_total}, particularly in view of the different physical scales of the measured regions inside different galaxies (see Section \\ref{measure}).  The adoption of various SFHs complements the interpretation of the data distributions with the model assuming simple stellar populations. However, it remains deficient for the current scenarios to well characterize all the data. As a number of studies have suggested, various dust grain properties and diverse spatial geometries are suspected of causing discrepancies in the relationship between attenuation and observed color indices at certain wavebands, which thus leads to different IRX-UV trends \\citep[e.g.,][]{2000ApJ...528..799W, 2006MNRAS.370..380I}. In Paper II, we will investigate the dependence of the IRX-UV relation on attenuation/extinction law and attempt to discover more parameters in the IRX-UV function, speci\ufb01cally in order to obtain solutions from an alternative viewpoint to the discrepancies between the observational data and the current scenarios.  Statistical studies of integrated galaxies have shown large dispersion in the IRX-UV relation \\citep[e.g.,][]{2005ApJ...619L..51B, 2007ApJ...655..863D, 2009ApJ...703..517D, 2007ApJS..173..185G}. Although stellar population age has been widely considered as the second parameter for introducing the deviation, until now there has not been sufficient evidence to confirm the age effect on the IRX-UV relation for integrated galaxies \\citep{2005ApJ...619L..55S, 2007ApJS..173..392J}. In this paper, the comparisons between the spatially resolved measurements and the integrated properties of galaxies provide implications about the conundrum in integrated studies of galaxies. The locations of integrated galaxies and galactic subregions in the IRX-UV diagrams (Figures \\ref{IRXUV_UV_M81}$-$\\ref{IRXUV_UV_NGC7331}) and the fractional contributions of component to integrated luminosities of galaxies (Table \\ref{contri}) indicate that, any single component lacks the ability to dominate the integrated luminosities of galaxies, and integrated measurements of normal galaxies are aggregations of various populations and different components within galaxies. As also implied in \\citet{2012A&A...539A.145B}, simple age tracers are not in a good position to estimate stellar population age in such a case. If galaxies exhibit prominent features in sub-structures at certain observational bands, the integrated measurements can have an effective impact. In most cases of integrated measurements which compromise different stellar populations in galaxies and consequently represent more complexities in SFH, contrary to spatially resolved studies, any significant variation in the IRX-UV trend as a function of stellar population age tends to become unclear and probably masked by other potential factors. Moreover, even though in the simple situation any other parameter is negligible, the effects of stellar population age are suggested to be more complicated than we have realized in the IRX-UV plane \\citep{2008MNRAS.386.1157C}."
    },
    "1208/1208.0444_arXiv.txt": {
        "abstract": "We present a derivation of the time evolution equations for the energy content of nonhelical magnetic fields and the accompanying turbulent flows from first principles of incompressible magnetohydrodynamics in the general framework of homogeneous and isotropic turbulence. This is then applied to the early Universe, i.e., the evolution of primordial magnetic fields. Numerically integrating the equations, we find that most of the energy is concentrated at an integral wavenumber scale $k_{I}$ where the turbulence turn over time equals the Hubble time. At larger length scales $L$, i.e., smaller wavenumbers $q = 2 \\pi / L\\ll k_I$, independent of the assumed turbulent flow power spectrum, mode-mode coupling tends to develop a small $q$ magnetic field tail with a Batchelor spectrum proportional to the fourth inverse power of $L$ and therefore a scaling for the magnetic field of $B \\sim L^{-5/2}$. ",
        "introduction": "The question of the origin and time evolution of primordial magnetic fields in the early Universe is an interesting and so far at best partially resolved problem in cosmology. It is possible that strong magnetic fields have been created on small scales in the early Universe by a cosmological process, for example at a phase transition \\cite{PhysRevD.55.4582} or, less likely, during inflation \\cite{PhysRevD.37.2743} (for an overview of possible magnetogenesis scenarios see, for example, Ref.~\\cite{PhysRep.348.163} or Ref.~\\cite{Subramanian:2009fu}). A central question then is whether these magnetic fields could evolve with time and transport some of their energy content to large scales to account for the recently claimed detection of intergalactic magnetic fields \\cite{PhysRevD.80.123012,Neronov02042010}.  Many attempts to study the evolution of primordial magnetic fields have been performed in the past \\cite{Dimopoulos:1996nq,Brandenburg:1996fc,Jedamzik:1996wp, Subramanian:1997gi,Son:1998my,Christensson:2000sp,Sigl:2002kt,Banerjee:2004df,Campanelli:2007tc,Kahniashvili:2010gp}. Numerical simulations are problematic as they lack the resolution required to give reliable predictions. A further complication in the study of the evolution of magnetic fields in the early Universe is the enormous cosmic expansion between a putative magnetogenesis scenario and the present, such that the smallest errors in extrapolation lead to large changes in the final prediction.  In this paper we take a semianalytic approach which derives the main time evolution equations from first principles of magnetohydrodynamics (MHD), employing some fairly generic assumptions. Our analysis follows a similar procedure as has been already applied to, for example, the solar wind \\cite{Grappin82,Grappin83} or the galactic dynamo problem \\cite{Kulsrud:1992rk}.  Two of us \\cite{PhysRevD.83.103005} have recently, for the first time, applied such an approach to the evolution of primordial magnetic fields. Considering only the most important large $q$ velocity-magnetic field mode-mode coupling source term to generate magnetic fields on small $q$ taken from Ref.~\\cite{Kulsrud:1992rk}, and building on the result of Banerjee and Jedamzik \\cite{Banerjee:2004df}, it has been established that cosmic expansion seems slow enough to allow for the generation of large-scale magnetic fields with a white noise, i.e., $B \\sim L^{-3/2}$, or even shallower spectrum, also in the absence of initial magnetic and velocity fields on such large scales \\cite{remark1}. However, the analysis is far from complete as the role of back-reaction of the turbulent medium onto the magnetic fields has not been accounted for. Here we present the full analysis including all source and sink terms for nonhelical fields.  This paper is structured as follows: In Sec.~\\ref{sec:TimeHomoIsoMHD} we give the evolution equations for the magnetic and turbulence fields in the general framework of homogeneous and isotropic noncompressible MHD. An outline of the lengthy derivation of these equations is delegated to Appendix \\ref{app:calc}, whereas a different form more suitable for numerical integrations is presented in Appendix \\ref{app:alternative}. In Sec.~\\ref{sec:EarlyUniverse} this is then extended to the situation in the early Universe with the results and conclusions from these considerations finally presented in Secs.~\\ref{sec:Results} and \\ref{sec:Conclusions}, respectively. ",
        "conclusions": "\\label{sec:Conclusions} Starting from first principles and reasonable assumptions about homogeneous and isotropic magneto-hydrodynamic turbulence, evolution equations for the spectral magnetic and kinetic energy densities have been derived. After adapting these equations to the use in an expanding Universe they were numerically integrated for a number of assumed initial conditions, as possibly resulting during an early magnetogenesis period. It has been found that, seemingly independent of initial conditions, a $B(L)\\sim L^{-5/2}$ tail on scales $L \\gg 2 \\pi/k_{I}$ always develops. This magnetic field spectrum may therefore be regarded as the natural one for nonhelical cosmic magnetic fields (cf. also to Ref.~\\cite{Durrer:2003ja}). At the same time most of the energy is concentrated on the integral scale which allows a rather generic prediction of coherence scale and strength of causally created primordial magnetic fields."
    },
    "1208/1208.5866_arXiv.txt": {
        "abstract": "{Both the absolute mass-loss rates and the mechanisms that drive the mass loss of late-type supergiants are still not well known. Binaries such as $\\alpha$~Sco provide the most detailed empirical information about the winds of these stars.} {The goal was to improve the binary technique for the determination of the mass-loss rate of \\object{$\\alpha$~Sco~A} by including a realistic density distribution and velocity field from hydrodynamic and plasma simulations.} {We performed 3D hydrodynamic simulations of the circumstellar envelope of $\\alpha$~Sco in combination with plasma simulations accounting for the heating, ionization, and excitation of the wind by the radiation of \\object{$\\alpha$~Sco~B}. These simulations served as the basis for an examination of circumstellar absorption lines in the spectrum of $\\alpha$~Sco~B as well as of emission lines from the \\object{Antares nebula}.} {The present model of the extended envelope of $\\alpha$~Sco reproduces some of the structures that were observed in the circumstellar absorption lines in the spectrum of $\\alpha$~Sco~B. Our theoretical density and velocity distributions of the outflow deviate considerably from a spherically expanding model, which was used in previous studies. This results in a higher mass-loss rate of $(2\\pm0.5)\\times10^{-6}\\ M_{\\odot}\\,\\mathrm{yr}^{-1}$. The hot \\ion{H}{ii} region around the secondary star induces an additional acceleration of the wind at large distances from the primary, which is seen in absorption lines of \\ion{Ti}{ii} and \\ion{Cr}{ii} at $-30\\ \\mathrm{km\\,s}^{-1}$.} {} ",
        "introduction": "\\label{introd} In the late stage of the stellar evolution stars suffer a substantial mass loss by a continuous wind. This phenomenon plays a key role for the final development of individual stars and for the enrichment of the interstellar medium. However, the physical mechanisms that govern the outflow of late-type giants and supergiants are not established. Possible driving mechanisms include Alfv\\'{e}n waves, acoustic waves, shock waves, and pulsations \\citep[see the reviews of][]{Lafon_Berruyer_1991, Willson_2000}. As it is hard to distinguish between different driving mechanisms from a theoretical point of view, there is still a need for empirical studies to provide constraints.  The most accurate mass-loss rates are obtained from observations of eclipsing binaries of $\\zeta$~Aur or VV~Cep type. These systems consist of a late-type supergiant and a main-sequence companion that can be used as a ``natural satellite'' to probe the circumstellar envelope. For the basic concepts of the binary technique see e.\\,g.\\ \\citet{Reimers_1987}, \\citet{Hack_Stickland_1987}, or \\citet{Baade_1998}. Observations with high spectral resolution of $\\zeta$~Aur \\citep{Baade_etal_1996} and $\\alpha$~Sco \\citep{Baade_Reimers_2007, Reimers_etal_2008} have shown that the circumstellar envelopes of these stars cannot be described by a continuous, spherically expanding model. One of the aims of this study was therefore to investigate the influence of the binary companion on the density and velocity distribution of the supergiant wind.  The $\\alpha$~Sco system (M1\\,Ib + B2.5\\,V) is of particular importance in this context, because its angular separation amounts to nearly 3\\arcsec, which allows the observation of the extended envelope with spatial resolution. \\citet{Reimers_etal_2008} presented the most recent extensive observations, which were carried out with the Ultraviolet and Visual Echelle Spectrograph (UVES) on the Kueyen telescope of the Very Large Telescope array (VLT) and provide spatially resolved emission spectra of the \\object{Antares} nebula. They showed that the mass-loss rate of $\\alpha$~Sco~A can be determined by measuring the spatial extent of the H$\\alpha$ emission. They determined a mass-loss rate of $\\dot{M}=(1.05\\pm0.3)\\times10^{-6}\\ M_\\odot\\,\\mathrm{yr}^{-1}$ as a best match between the measured H$\\alpha$ distribution and plasma simulations, which were based on a spherically expanding model of the circumstellar envelope and led to a nearly spherical \\ion{H}{ii} region around the secondary that is rotationally symmetric with respect to the line connecting the two stars. They reported an apparent increase of the mass-loss rate towards larger distances from the supergiant. Moreover, they noted that the observed [\\ion{Fe}{ii}] emission cannot be explained in the framework of their simplified model.  Mid-infrared observations at the Keck~\\textsc{ii} telescope of the dust in the circumstellar envelope of $\\alpha$~Sco indicate a non-uniform dust distribution that appears to be related to discrete ejections of mass from $\\alpha$~Sco~A \\citep{Marsh_etal_2001}. This is consistent with interferometric observations that were carried out at the William Herschel Telescope by \\citet{Tuthill_etal_1997} and revealed asymmetric structures (``hotspots'') at the surface of the supergiant.  \\citet{Baade_Reimers_2007} presented spectra of $\\alpha$~Sco~B, which were obtained with the Goddard High-Resolution Spectrograph (GHRS) on the Hubble Space Telescope (HST), including a number of circumstellar absorption lines. These absorption lines exhibit multi-component structures that are not compatible with a continuous wind, and the derived column densities indicated a complex differential depletion as a result of dust formation. They showed that these structures can be explained in principle either by a model consisting of concentric shells with different densities, consistent with a highly variable mass-loss rate, or by discrete clumps in the line of sight due to unknown ejection processes.  The goal of this work was to construct a more realistic model of the $\\alpha$~Sco system, including hydrodynamic and plasma effects, in order to gain quantitative insight into the dynamics of the circumstellar envelope. After a description of our hydrodynamic and plasma simulations of the $\\alpha$~Sco system and its \\ion{H}{ii} region in Sect.~\\ref{modcal}, we present a comparison to observations in Sect.~\\ref{result}. ",
        "conclusions": "The results of the combination of observations with hydrodynamic and plasma simulations presented in this work show that the circumstellar envelope of the $\\alpha$~Sco system is strongly influenced by dynamic effects. A calculation of absorption line profiles based on the simulated density and velocity distributions and ionization structure, and a comparison to GHRS/HST and UVES/VLT spectra reveal that the mass-loss rate was underestimated by a factor of two in earlier studies that were based on an undisturbed, spherically symmetric circumstellar shell. The resulting mass-loss rate of $\\dot{M}=2\\times10^{-6}\\ M_\\odot\\,\\mathrm{yr}^{-1}$ is confirmed by an analysis of the H$\\alpha$ emission from the Antares nebula, which is based on spatially resolved emission distributions observed with UVES.  The theoretical absorption line profiles exhibit a multi-component structure as a natural result of the influence of the hot \\ion{H}{ii} region that is moving with the B star through the wind of the primary. Three of the four observed components can be explained accordingly. However, the origin of the additional component at $-14\\ \\mathrm{km\\,s}^{-1}$ seen in most absorption lines (Fig.~\\ref{obsabs}) remains uncertain and is probably an indication of time-dependent mass-loss.  The structure of the [\\ion{Fe}{ii}] line emission of the Antares nebula as observed with UVES cannot be reproduced correctly. A more sophisticated model including exact radiative transfer calculations and time-dependent simulations of the ionization balance allowing for advection effects would help to understand the remaining discrepancies between theory and observations. In this context, it would be desirable to parallelize the hydro code, which will make it possible to calculate a more extensive grid of models. Future projects may deal with these improvements of the model calculations.  A number of empirical mass-loss rates of $\\zeta$~Aur systems were determined on the basis of spherically symmetric density and velocity distributions \\citep[see e.\\,g.][]{Che_etal_1983, Baade_etal_1996}. Some of these systems, e.\\,g.\\ $\\zeta$~Aur and 31~Cyg, also contain \\ion{H}{ii} regions, which may lead to even more severe dynamic effects than in $\\alpha$~Sco due to the much smaller orbital periods. The current values of the mass-loss rates of these stars may have to be revised on the basis of more realistic models including dynamic processes as presented in this work for $\\alpha$~Sco."
    },
    "1208/1208.0358_arXiv.txt": {
        "abstract": "We compare the current LIGO/Virgo upper limits on double compact object volumetric merger rates with our theoretical predictions. Our optimistic models are a factor of $\\sim 3$ below the existing upper limits for massive BH-BH systems with total mass $50$--$70\\msun$, suggesting that a small increase in observational sensitivity may bring the first detections. The LIGO/Virgo gravitational wave observatories are currently being upgraded to advanced design sensitivity. If a sizeable population of BH-BH binaries is detected, the maximum total binary mass of this population will discriminate between two general families of common envelope models. If no binaries are detected, the new upper limits will provide astrophysically useful information about the environment and physical processes (e.g., metallicity of host galaxies or BH natal kicks) crucial to the formation of binaries containing black holes. For NS-NS systems, our predicted rates are $\\sim 3$ orders of magnitude below the current upper limits; even if advanced instruments reach their design sensitivities (factors of $10\\times$ in distance, and $1,000\\times$ in volumetric rate) the detection of NS-NS systems is not assured. However, we note that although our predicted NS-NS merger rates are consistent with estimates derived from Galactic NS-NS binaries and short GRBs, they are on the low side of these empirical estimates. ",
        "introduction": "Most massive stars are found in binary systems (Goodwin et al. 2007). During the evolution of these stars the binaries can experience component merger during common envelope (CE) phases (Webbink 1984) or disruption during supernova (SN) explosions (Tauris \\& Takens 1998) in which individual stars form neutron stars (NSs) or black holes (BHs). The massive binaries which survive these processes form double compact objects: NS-NS, BH-BH, or mixed BH-NS systems. These remnant systems are subsequently subject to angular momentum loss via the emission of gravitational waves (GWs) and their orbital separation decreases (Peters \\& Mathews 1963; Weisberg \\& Taylor 2005). Finally, the two compact objects collide and merge into a single compact object giving rise to a strong GW signal (Einstein 1918). The LIGO/Virgo network of ground-based interferometric observatories has been designed to search for these signals.\\footnote{http://www.ligo.caltech.edu/; https://wwwcascina.virgo.infn.it/}  Initial LIGO/Virgo observations were concluded in $2012$ without the detection of a GW signal. The instruments are currently being upgraded and the network will resume its operation in several years with advanced sensitivity. Various predictions for near-future detection chances were compiled and presented by Abadie et al. (2010).  One of the most promising sources for these advanced GW detectors is the inspiral and merger of binary NS systems. There are several known Galactic double neutron star binaries with merger times shorter than the Hubble time (e.g., Kim, Kalogera \\& Lorimer 2010). Double black hole binaries (BH-BH), on the other hand, remain undetected. Recent theoretical predictions indicate that these systems may either dominate GW observations (e.g., Belczynski et al. 2010a) or be totally absent in local Universe (e.g., Mennekens \\& Vanbeveren 2013).  In this study we compare the rates from our evolutionary calculations of double compact object mergers  (Dominik et al. 2012) with the existing LIGO/Virgo upper limits (Abadie et al. 2012; Aasi et al. 2013). We then make predictions for the science that will be probed with future upper limits and/or detections of double compact object mergers with advanced GW instruments. ",
        "conclusions": "Our results show that a small improvement in LIGO/Virgo's current sensitivity to double compact object mergers, by a factor of $\\sqrt[3]{3}$ in distance (factor of 3 in volume), will either bring the first detections of massive BH-BH binaries ($M_{\\rm tot}>50\\msun$) or through non-detection exclude entire families of evolutionary models (see the proximity of Model A to the LIGO/Virgo upper limits in Figure~\\ref{std}).  If discoveries of massive BH-BH binaries are made, even at low rates, it will indicate that stars just beyond the main sequence (i.e., on the Hertzsprung gap) survive common envelope and form close double compact objects (see Figures~\\ref{std}, \\ref{metals}, \\ref{kicks}, and differences between models A and B).  If upper limits are improved by a factor of $100$ and we remain without the detection of BH-BH binaries, it will indicate that other factors than common envelope are responsible for eliminating the BH-BH formation channels. Our analysis indicates that high black hole natal kicks and high metallicity stellar environments are two crucial factors that significantly reduce the BH-BH merger rates (see Figures~\\ref{metals}, \\ref{kicks} and the associated merger rate reductions). These two factors show a degeneracy when it comes to the rate of BH-BH mergers, so additional constraints are required to distinguish their relative importance.  If upper limits are improved by factor of $1,000$ and no BH-BH discoveries are made, then the above degeneracy is lifted and we can conclude that BHs receive high natal kicks {\\em and}\\/ that the stellar populations within reach of LIGO are typically of high (solar-like) metallicity (see Figure~\\ref{kicks2}, where the BH-BH rates fall below the enhanced upper limits). This is demonstrated by our V15 model in which we have increased the magnitude of BH natal kicks and the high metallicity content of the local Universe to the largest values with reasonable limits: it is difficult to imagine BHs receiving higher kicks than NSs (see Figure~\\ref{bhkick}), and less than $10\\%$ of stars in the local Universe having sub-solar metallicity (e.g., Panter et al. 2008).  The above conclusions apply only within the framework of our evolutionary model. Our model successfully recovers the observed BH mass spectrum both in the Galaxy and in other environments (Belczynski et al. 2010b). It also provides a physical explanation for the existence of the mass gap between neutron stars and black holes, failing to produce compact objects in the mass range $2$--$5\\msun$ (Belczynski et al. 2012). However, our model is unable to reproduce the companion mass distribution observed in Galactic BH transients. The observed distribution peaks at $\\sim 0.6\\msun$ while most of our predicted distributions peak at $\\sim 1\\msun$. This discrepancy may arise from our poor understanding of magnetic braking in these low mass companions along with uncertainties in low mass stellar models. It is also possible that the discrepancy arises from observational biases (Wiktorowicz, Belczynski, \\& Maccarone 2013). These factors should not play a significant role in the modeling of massive stars and the formation of NS-NS, BH-NS, or BH-BH binaries.  A different reduction factor in BH-BH formation was proposed by  Mennekens \\& Vanbeveren (2013). These authors claim that strong Luminous Blue Variable (LBV) winds may inhibit formation of close BH-BH binaries. They argue that the action of these very intense winds from BH progenitors leads to an increase in orbital separation such that BH-BH binaries with merger times below a Hubble time no longer form. This finding is hard to reconcile with observations, as a number of Galactic and extra-galactic binaries with black holes on close orbits are known (e.g., extra-galactic IC10 X-1 with $P_{\\rm orbital}=33\\,$h or Galactic GRO J0422+32 with $P_{\\rm orbital}=5.1\\,$h both hosting massive $\\gtrsim10\\msun$ BHs).  If NS-NS mergers are found before the advanced instruments reach their design sensitivity (a improvement by a factor of $\\sim1,000$ in volume) it will indicate that our predicted NS-NS merger rates densities are too low. The astrophysical implications of such a potential finding are now under study (Dominik et al., in preparation). Alternatively, it is worth noting that even with an increase in the existing upper limits by factor of $\\sim1,000$ some of our model predictions, as well as low values of rates derived from observations of Galactic NS-NS binaries and cosmic short GRBs, allow for the possibility of non-detection of NS-NS mergers (see Figure~\\ref{kicks2}; some of our models, as well as empirical estimates, allow for lower merger rate densities than those shown).  In conclusion, there exist a wide range of possibilities for the rates of compact binary mergers in advanced detectors like LIGO and Virgo. The first detections may be made with small improvements over the instruments sensitivity. Alternatively, even at full advanced design sensitivity there may be no detections of double compact object mergers. Regardless of whether there are detections or more stringent upper limits on the rates, LIGO/Virgo will address a number of important science problems in our understanding of compact binary formation and evolution."
    },
    "1208/1208.3737_arXiv.txt": {
        "abstract": "We present the ``Very Large Array survey of Advanced Camera for Surveys Nearby Galaxy Survey Treasury galaxies (VLA-ANGST).'' VLA-ANGST is a National Radio Astronomy Observatory Large Program consisting of high spectral ($0.6-2.6$\\,\\kms) and spatial ($\\sim 6\\arcsec$) resolution observations of neutral, atomic hydrogen (\\hi) emission toward 35 nearby dwarf galaxies from the ANGST survey. ANGST is a systematic Hubble Space Telescope survey to establish a legacy of uniform multi-color photometry of resolved stars for a volume-limited sample of nearby galaxies (D $\\lesssim$ 4 Mpc). VLA-ANGST provides VLA \\hi\\ observations of the sub-sample of ANGST galaxies with recent star formation that are observable from the northern hemisphere and that were not observed in the ``The \\hi\\ Nearby Galaxy Survey'' (THINGS). The overarching scientific goal of VLA-ANGST is to investigate fundamental characteristics of the neutral interstellar medium (ISM) of dwarf galaxies. Here we describe the VLA observations, the data reduction, and the final VLA-ANGST data products. We present an atlas of the integrated \\hi\\ maps, the intensity-weighted velocity fields, the second moment maps as a measure for the velocity dispersion of the \\hi, individual channel maps, and integrated \\hi\\ spectra for each VLA-ANGST galaxy. We closely follow the observational setup and data reduction of THINGS to achieve comparable sensitivity and angular resolution. A major difference between VLA-ANGST and THINGS, however, is the high velocity resolution of the VLA-ANGST observations (0.65 and 1.3\\,\\kms\\ for the majority of the galaxies). The VLA-ANGST data products are made publicly available through a dedicated webpage\\footnote{\\url{https://science.nrao.edu/science/surveys/vla-angst}}. With available star formation histories from resolved stellar populations and lower resolution ancillary observations from the far infrared to the ultraviolet, VLA-ANGST will enable detailed studies of the relationship between the ISM and star formation in dwarf galaxies on a $\\sim 100$\\,pc scale. ",
        "introduction": "\\label{sec:intro}  Star formation is driven by complicated interactions between gas and stars.  Untangling the interplay between these processes is difficult, because in most cases, the events that trigger star formation are not obvious, nor are those that shape the structure and dynamics of the ISM. For a full understanding it is therefore indispensable to obtain a comprehensive view of all processes that come together to form stars, stellar associations, and stellar clusters.  Observationally, one requires knowledge of the gas distribution and kinematics as well as the stellar energy input into the ISM over time.  In recent years, large systematic surveys have made superb progress on the first of these requirements.  The number of nearby galaxies for which high-quality \\hi\\ data is available has dramatically increased in the last few years, including campaigns such as THINGS \\citep[``The \\ion{H}{1} Nearby Galaxy Survey'';][]{wal08}, FIGGS \\citep[``Faint Irregular Galaxies GMRT Survey'';][]{beg08}, LITTLE-THINGS \\citep[``LITTLE: Local Irregulars That Trace Luminosity Extremes'';][]{hun07}, SHIELD \\citep[``Survey of \\hi\\ in Extremely Low-mass Dwarfs'';][]{can11a}, LVHIS \\citep[``The Local Volume \\ion{H}{1} Survey'';][]{kor08}, WHISP \\citep[``Westerbork observations of neutral Hydrogen in Irregular and Spiral galaxies'';][]{vdh01}, and HALOGAS \\citep[``The Westerbork Hydrogen Accretion in Local Galaxies'';][]{hea11}. The difficult work of mapping the molecular medium in the brighter galaxies has begun as well, e.g. in BIMA SONG \\citep[``The BIMA Survey of Nearby Galaxies''][]{hel03}, HERACLES, \\citep[``The HERA CO Line Extragalactic Survey'';][]{ler09}, and STING \\citep[``CARMA Survey Toward Infrared-bright Nearby Galaxies'';][]{rah12}.  Unfortunately, the needed measurements of time-resolved stellar energy input are more difficult to acquire.  Large systematic surveys in the far-ultraviolet (e.g., obtained with the {\\it GALEX} telescope) and the far-infrared (e.g., the Local Volume Legacy survey \\citep[LVL;][]{dale09}, SINGS \\citep[``Spitzer Infrared Nearby Galaxies Survey'';][]{ken03}, Kingfish \\citep[``Key Insights on Nearby Galaxies: a Far-Infrared Survey with Herschel'';] []{ken11}, etc) have made excellent progress in measuring the recent star formation rate (SFR), while surveys like 11HUGS \\citep[``The 11\\,Mpc H$\\alpha$ UV Galaxy Survey'';][]{lee04} have provided the H$\\alpha$ mapping needed to trace star formation (SF) on much shorter timescales.  However, these approaches to measuring star formation lack all but the broadest time resolution, with the different tracers being sensitive to SF over different timescales. H$\\alpha$ is emitted on the timescale of O-star lifetimes ($\\lesssim 5$\\,Myr) and the far-ultraviolet on the timescale of A-star lifetimes ($\\lesssim 100$\\,Myr). The far-infrared is sensitive to timescales similar to the $\\lesssim 100$\\,Myr of far-ultraviolet heating. These timescales are not necessarily well-matched to the relevant energy input timescales for the gas.  For the survey presented here, we take a different approach, and focus \\hi\\ observations on galaxies which are sufficiently close that their stellar populations can be resolved with the {\\it Hubble Space Telescope} (HST).  The resulting color--magnitude diagrams (CMDs) allow one to construct spatially-resolved star formation histories (SFHs) via analyses of galaxies' stellar populations, and thus reveal the time-resolved SFR of these nearby galaxies over timescales of several hundred Myr at $\\sim5-10\\arcsec$ spatial resolution \\citep[e.g.,][]{doh02,can03,wei08,wil10,mcq10,crn11,can11b,wei11}. With nearly 300 orbits of {\\it HST} time, one of the most ambitious programs to obtain spatially-resolved SFHs is the ACS Nearby Galaxy Survey Treasury program \\citep[ANGST;][]{dal09}. The ANGST observations enable to map spatially-resolved SFHs for a volume-limited sample of 69 nearby ($<4$\\,Mpc) galaxies, probing both group and field environments. These data provide an entirely new, statistically significant view on the SFH of the local universe dwarf galaxy population.  The survey presented here, VLA-ANGST (``Very Large Array survey of ACS Nearby Galaxy Survey Treasury galaxies''), complements ANGST with high spatial and spectral resolution data cubes of the atomic gas traced by the 21\\,cm line of neutral atomic hydrogen (\\hi). VLA-ANGST was designed to aim for the best available resolution and sensitivity, using the NRAO Very Large Array (VLA) in multiple configurations in a Large Program worth $\\sim 500$\\,hours of observing time to achieve that goal. VLA-ANGST is designed to match the \\hi\\ spatial resolution ($\\sim 6\\arcsec$) to the cell sizes over which the SFHs can be determined. Furthermore, the majority of the VLA-ANGST galaxies were observed at very high spectral resolution of $0.6-1.3$\\,\\kms\\ which is important for detailed ISM dynamic modeling studies of the rather low-mass objects which dominate the galaxy population within $4$\\,Mpc.  The VLA-ANGST survey has a number of features that make it a valuable addition to the many other \\hi\\ surveys, beyond the existence of resolved stellar population data.  First, the galaxies in VLA-ANGST are all quite close, which ensures high linear resolution for studying small-scale features in the \\hi\\ distribution.  Second, because care was taken to match the observational setups of THINGS, the VLA-ANGST survey can be readily combined with the THINGS survey, giving much broader coverage towards low galaxy masses.  When further combined with surveys like LITTLE-THINGS and SHIELD, which have used a similar observational strategy, we will have a relatively uniform database of \\ion{H}{1} toward $>100$ objects spanning a large variety of galaxy types that is comparable in terms of sensitivity, angular, and spectral resolution.   The variety of galaxy types in the sample allows the study of (1) the response of gas and star formation to the propagation of spiral arms and to interactions, as seen in massive spirals and starburst galaxies (2) star formation propagation in the absence of strong perturbations of the gas density; gas rich dwarves are ideal for such a study due to their lack of internal shear and spiral density waves (3) star formation triggered in unusual kinematic environments such as in tidal dwarfs, and (4) dIrr/dSph transition-type galaxies, whose lack of current SF but ample gas reservoirs allow studies of galaxies that possess the raw material for star formation, but somehow remain dormant.   In the following we present the data of the VLA-ANGST survey, and in-depth scientific analyses will follow in subsequent publications. First analyses based on VLA-ANGST data cubes are provided by \\citet{war11} on the energy requirements of large \\hi\\ holes, by \\citet{war12} on the distribution of cold and warm \\hi, and by \\citet{sti12} on the global \\hi\\ velocity dispersion profiles correlated to the properties of the galaxies. Section\\,\\ref{sec:sample} describes the selection of the targets, followed by the observational setup and data reduction (Section\\,\\ref{sec:reduction}). Our data products are presented in Section\\,\\ref{sec:dataProducts} and a summary of the VLA-ANGST given in Section\\,\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  Here we present the sample selection criteria, observational parameters, data reduction procedures, and data product description of the VLA-ANGST survey, a Large VLA project that targets nearby, mostly dwarf irregular galaxies. Of the 35 galaxies in the survey, we detect \\hi\\ in 29. The calibrated VLA data for these objects are publicly available at \\url{https://science.nrao.edu/science/surveys/vla-angst}. This leads to the following data products:   \\begin{itemize}  \\item Global \\hi\\ spectra for all galaxies, derived from the masked, flux-corrected, primary beam corrected, natural-weighted data cubes.  \\item \\hi\\ data cubes of both natural and robust weighting. The cubes are not primary beam or flux corrected and they are unmasked.  \\item Integrated intensity maps (moment 0 maps) in units of Jy\\,beam$^{-1}$\\,\\kms\\ as well as converted to \\hi\\ column densities. These maps were derived from the masked, flux-corrected, primary beam corrected data cubes. We offer both, natural and robust-weighted maps for download.  \\item The intensity-weighted velocity field maps (moment 1 maps), derived from the same data products as the integrated intensity maps.  \\item Second moment maps which give a measure for the velocity dispersion of the gas; derived from the same data products as the integrated intensity maps.  \\end{itemize}  This paper presents the observations; scientific analyses will follow in subsequent publications. A study that compares the \\hi\\ kinematics of large \\hi\\ shells to the supernovae and stellar wind output of the underlying stellar populations is presented in \\citet{war11}. Detection and characterization of narrow \\hi\\ components that presumably trace cold \\hi\\ are provided in \\citet{war12}. \\citet{sti12} correlate averaged \\hi\\ dispersion values to the physical properties of the host galaxies. The true value of our \\hi\\ survey is further unlocked by the extensive, multi-wavelength ancillary data that is available for many of our objects. The VLA-ANGST, THINGS, LITTLE THINGS, and SHIELD data products are furthermore similar in sensitivity, spatial and spectral resolution and provide \\hi\\ data cubes for $>100$ galaxies.     \\begin{deluxetable}{llrrrrrrrc}  \\tablewidth{0pt} \\tabletypesize{\\tiny}% \\rotate \\tablecaption{General properties of the VLA-ANGST galaxy sample.\\label{tab:sample}} \\tablehead{ \\colhead{(1)}           & \\colhead{(2)}      & \\colhead{(3)}          & \\colhead{(4)}  & \\colhead{(5)}    & \\colhead{(6)}  & \\colhead{(7)}  & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)}\\\\  \\colhead{Name}           & \\colhead{Alt. Name}      & \\colhead{RA}          & \\colhead{DEC}  & \\colhead{D}    & \\colhead{$D_{\\rm 25}$}  & \\colhead{$M_{\\rm B}$}  & \\colhead{$\\nu L_{\\nu}$} & \\colhead{Type}& \\colhead{$SFR$}\\\\  \\colhead{}           & \\colhead{}      & \\colhead{(J2000)}          & \\colhead{(J2000)}  & \\colhead{}    & \\colhead{}  & \\colhead{}  & \\colhead{($3.6$\\,$\\mu$m)} & \\colhead{}& \\colhead{(UV)}\\\\  \\colhead{}           & \\colhead{}      & \\colhead{[$^{h}$:$^{m}$:$^{s}$]}          & \\colhead{[$\\degr$:$\\arcmin$:$\\arcsec$]}  & \\colhead{[Mpc]}    & \\colhead{[kpc]}  & \\colhead{[mag]}  & \\colhead{[$10^{6}$\\,L$_{\\sun}$]} & \\colhead{}& \\colhead{[$10^{-3}$\\,\\sfr]}  }  \\startdata NGC\\,247              & ESO\\,540-G022, UGCA\\,11      & 00:47:08.3 & -20:45:36 & 3.50 & 15.7   & -17.86  &  270.52   & 7   &  229  \\\\ DDO\\,6                & ESO\\,540-G031, UGCA\\,15      & 00:49:49.3 & -21:00:58 & 3.31 & 1.6    & -12.41  &  \\nodata  & 10  &  1.4    \\\\ NGC\\,404              & UGC\\,718                     & 01:09:26.9 & 35:43:03  & 3.05 & 2.2    & -16.21  &  \\nodata  & -1  &  6.6    \\\\ KKH\\,37               & LEDA\\,95597                  & 06:47:45.8 & 80:07:26  & 3.26 & 1.1\\tablenotemark{a}    & -11.17  &  0.69     & 10  &  0.3    \\\\ UGC\\,4483             & CGCG\\,331-051                & 08:37:03.0 & 69:46:31  & 3.41 & 1.2    & -12.71  &  0.57     & 10  &  4.0    \\\\ KK\\,77                & LEDA\\,166101                 & 09:50:10.0 & 67:30:24  & 3.55 & 2.5\\tablenotemark{a}    & -11.45  &  \\nodata  & -3  &  \\nodata \\\\ BK3N                  & PGC\\,28529                   & 09:53:48.5 & 68:58:09  & 3.86 & 0.6    & -9.15   &  $<$0.03  & 10  &  0.3    \\\\ AO\\,0952+69\\tablenotemark{b}           & Arp's Loop                   & 09:57:29.0 & 69:16:20  & 3.78 & 2.0    & -11.09  &  \\nodata  & 10  &  1.1    \\\\ Sextans\\,B            & UGC\\,5373, DDO\\,70           & 10:00:00.1 & 05:19:56  & 1.39 & 2.1    & -13.87  &  2.48     & 10  &  4.5    \\\\ NGC\\,3109             & ESO\\,499-G036, DDO\\,236      & 10:03:07.2 & -26:09:36 & 1.26 & 6.2    & -15.11  &  12.40    & 10  &  28.7  \\\\ Antlia                & PGC\\,29194                   & 10:04:04.0 & -27:19:55 & 1.29 & 0.8    & -9.36   &  \\nodata  & 10  &  1.8    \\\\ KDG\\,63               & UGC\\,5428\\,DDO71             & 10:05:07.3 & 66:33:18  & 3.53 & 1.7    & -11.73  &  1.36     & -3  &  \\nodata \\\\ Sextans\\,A            & UGCA 205, DDO 075            & 10:11:00.8 & -04:41:34 & 1.38 & 2.2    & -13.84  &  1.82     & 10  &  12.3    \\\\ HS\\,117               & \\nodata                      & 10:21:25.2 & 71:06:58  & 3.82 & 1.7    & -11.41  &  1.08     & 10  &  \\nodata \\\\ DDO\\,82               & UGC\\,5692                    & 10:30:35.0 & 70:37:10  & 3.80 & 3.8    & -14.33  &  \\nodata  & 9   &  2.6    \\\\ KDG\\,73               & PGC\\,32667                   & 10:52:55.3 & 69:32:45  & 4.03 & 0.7    & -10.94  &  0.81     & 10  &  0.3    \\\\ NGC\\,3741             & UGC\\,6572                    & 11:36:06.4 & 45:17:07  & 3.24 & 1.9    & -13.17  &  1.29     & 10  &  6.2    \\\\ DDO\\,99               & UGC\\,6817                    & 11:50:53.0 & 38:52:50  & 2.59 & 3.1    & -13.37  &  1.43     & 10  &  5.5    \\\\ NGC\\,4163             & NGC\\,4167, UGC\\,7199         & 12:12:08.9 & 36:10:10  & 2.86 & 1.6    & -13.65  &  3.66     & 10  &  4.0    \\\\ NGC\\,4190             & UGC\\,07232                   & 12:13:44.6 & 36:38:00  & 3.50\\tablenotemark{c} & 1.7    & -14.20  &  5.80     & 10  &  10.1    \\\\ DDO\\,113              & UGCA\\,276                    & 12:14:57.9 & 36:13:08  & 2.95 & 1.3    & -11.65  &  0.61     & 10  &  \\nodata \\\\ MCG\\,+09-20-131       & CGCG\\,269-049                & 12:15:46.7 & 52:23:15  & 1.60\\tablenotemark{d} & 0.6    & -10.72  &  \\nodata  & 10  &  0.4   \\\\ DDO\\,125              & UGC\\,7577                    & 12:27:41.8 & 43:29:38  & 2.58 & 3.2    & -14.11  &  4.15     & 10  &  5.3    \\\\ UGCA\\,292             & PGC\\,42275                   & 12:38:40.0 & 32:46:00  & 3.62 & 1.1    & -11.69  &  0.48     & 10  &  2.9    \\\\ GR\\,8                 & UGC\\,8091, DDO\\,155          & 12:58:40.4 & 14:13:03  & 2.08 & 0.7    & -11.98  &  0.34     & 10  &  2.4    \\\\ UGC\\,8508             & I\\,Zw\\,60                    & 13:30:44.4 & 54:54:36  & 2.58 & 1.3    & -12.94  &  1.40     & 10  &  \\nodata \\\\ DDO\\,181              & UGC\\,8651                    & 13:39:53.8 & 40:44:21  & 3.14 & 2.1    & -13.03  &  1.51     & 10  &  3.8    \\\\ DDO\\,183              & UGC\\,8760                    & 13:50:51.1 & 38:01:16  & 3.22 & 2.1    & -13.09  &  1.63     & 9   &  3.2    \\\\ KKH\\,86               & LEDA\\,2807150                & 13:54:33.6 & 04:14:35  & 2.59 & 0.5\\tablenotemark{a}    & -10.19  &  0.22     & 10  &  0.1    \\\\ UGC\\,8833             & PGC\\,49452                   & 13:54:48.7 & 35:50:15  & 3.08 & 0.8    & -12.29  &  0.60     & 10  &  1.4    \\\\ KK\\,230               & KKR\\,3                       & 14:07:10.7 & 35:03:37  & 1.97 & 0.3\\tablenotemark{a}    & -8.57   &  0.05     & 10  &  0.2    \\\\ DDO\\,187              & UGC\\,9128                    & 14:15:56.5 & 23:03:19  & 2.21 & 1.1    & -12.34  &  0.39     & 10  &  1.1    \\\\ DDO\\,190              & UGC\\,9240                    & 14:24:43.5 & 44:31:33  & 2.79 & 1.5    & -14.13  &  3.12     & 10  &  6.1    \\\\ KKR\\,25               & LEDA\\,2801026                & 16:13:47.6 & 54:22:16  & 1.93 & 0.6\\tablenotemark{a}    & -9.98   &  $<$0.02  & 10  &  \\nodata \\\\ KKH\\,98               & LEDA\\,2807157                & 23:45:34.0 & 38:43:04  & 2.54 & 0.8\\tablenotemark{a}    & -10.32  &  0.40     & 10  &  0.6    \\\\ \\enddata  \\tablerefs{(5) tip of the red giant branch distances from \\citet{dal09}; (6) taken from \\citet{dal09} and converted to physical diameters; (7) apparent blue magnitudes from \\citet{kar04} and converted to absolute blue magnitudes; (8) converted from infrared fluxes given by \\citet{dale09}; (9) T-type form \\citet{dal09}; (10) converted from GALEX FUV asymptotic magnitudes given by \\citet{lee11} and using $SFR=1.4\\times10^{-28}\\,L_{\\nu}$(erg\\,s$^{-1}$\\,Hz$^{-1}$)  \\citep{ken98}}  \\tablenotetext{a}{for the KK-listed objects, the diameters are taken at a 26.5 mag\\,arcsec$^{-2}$ surface brightness level} \\tablenotetext{b}{object might be a feature in the spiral arm of M\\,81 rather than a galaxy} \\tablenotetext{c}{TRGB distance from \\citet{kar04}} \\tablenotetext{d}{the TRGB branch was not unambiguously identified in \\citet{dal09}}  \\end{deluxetable}    \\clearpage \\begin{deluxetable}{lcccccccccccccc} \\tablewidth{0pt} \\tabletypesize{\\scriptsize} \\rotate  \\tablecaption{List of Observations \\label{tab:obs}} \\tablehead{ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} & \\colhead{(10)} & \\colhead{(11)} & \\colhead{(12)} & \\colhead{(13)} & \\colhead{(14)} & \\colhead{(15)} \\\\   \\colhead{Galaxy} & \\colhead{Conf.} & \\colhead{Project} & \\colhead{Date} & \\colhead{RA} & \\colhead{Dec} & \\colhead{Equ.} & \\colhead{Cal} & \\colhead{Mode} & \\colhead{BW} & \\colhead{Chan} & \\colhead{$\\Delta$ v} & \\colhead{$\\nu_{\\mathrm{obs}1}$} & \\colhead{v$_{\\mathrm{obs}1}$} & \\colhead{N$_{EVLA}$} \\\\   \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{[yyyy-mm-dd]} & \\colhead{[$^{h}$:$^{m}$:$^{s}$]} & \\colhead{[$\\degr$:$\\arcmin$:$\\arcsec$]} & \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{[MHz]} & \\colhead{\\#} & \\colhead{[km s$^{-1}$]} & \\colhead{[MHz]} & \\colhead{[km s$^{-1}$]} & \\colhead{}  }  \\startdata  NGC\\,247        & BnA & AO215 & 2007-10-10  & 00:47:08.5 & -20:45:37 & 2000 & 0110-076 & 4   & 1.56 & 128 &  2.6 & 1419.098\\tablenotemark{a} & \\nodata & 12 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.222\\tablenotemark{a} & \\nodata &    \\\\ NGC\\,247        & BnA & AO215 & 2007-10-11  & 00:47:08.5 & -20:45:37 & 2000 & 0110-076 & 4   & 1.56 & 128 &  2.6 & 1419.098\\tablenotemark{a} & \\nodata & 12 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.222\\tablenotemark{a} & \\nodata &    \\\\ NGC\\,247        & B   & AO215 & 2008-01-11  & 00:47:08.5 & -20:45:37 & 2000 & 0110-076 & 4   & 1.56 & 128 &  2.6 & 1419.008\\tablenotemark{a} & \\nodata & 13 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.132\\tablenotemark{a} & \\nodata &    \\\\ NGC\\,247        & B   & AO215 & 2008-01-12  & 00:47:08.5 & -20:45:37 & 2000 & 0110-076 & 4   & 1.56 & 128 &  2.6 & 1419.008\\tablenotemark{a} & \\nodata & 13 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.132\\tablenotemark{a} & \\nodata &    \\\\ NGC\\,247        & CnB & AO215 & 2008-02-17  & 00:47:08.5 & -20:45:37 & 2000 & 0110-076 & 4   & 1.56 & 128 &  2.6 & 1419.059\\tablenotemark{a} & \\nodata & 13 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.182\\tablenotemark{a} & \\nodata &    \\\\ NGC\\,247        & DnC & AO215 & 2008-06-12  & 00:47:08.5 & -20:45:37 & 2000 & 0116-208 & 4   & 1.56 & 128 &  2.6 & 1419.233\\tablenotemark{a} & \\nodata & 15 \\\\ \\\\ DDO\\,6          & BnA & AO215 & 2007-10-05  & 00:49:49.2 & -21:00:54 & 2000 & 0145-275 & 2AD & 0.78 & 256 & 0.6  & 1419.010                  & \\nodata & 12 \\\\ DDO\\,6          & BnA & AO215 & 2007-10-07  & 00:49:49.2 & -21:00:54 & 2000 & 0145-275 & 2AD & 0.78 & 256 & 0.6  & 1419.010                  & \\nodata & 12 \\\\ DDO\\,6          & CnB & AO215 & 2008-02-16  & 00:49:49.2 & -21:00:54 & 2000 & 0145-275 & 2AC & 0.78 & 256 & 0.6  & 1418.953                  & \\nodata & 13 \\\\ DDO\\,6          & DnC & AO215 & 2008-06-12  & 00:49:49.2 & -21:00:54 & 2000 & 0116-208 & 2AC & 0.78 & 256 & 0.6  & 1419.158                  & \\nodata & 15 \\\\ DDO\\,6          & DnC & AO215 & 2008-07-11  & 00:49:49.2 & -21:00:54 & 2000 & 0116-208 & 2AC & 0.78 & 256 & 0.6  & 1419.140                  & \\nodata & 16 \\\\ \\\\ NGC\\,404        & B   & AO215 & 2007-11-13  & 01:09:27.0 & +35:43:04 & 2000 & 0119+321 & 2AD & 0.78 & 256 & 0.6  & 1420.598\\tablenotemark{a} & \\nodata & 12 \\\\ NGC\\,404        & C   & AC459 & 1996-01-01  & 01:06:39.0 & +35:28:00 & 1950 & 0116+319 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & -56.0   &  0 \\\\ NGC\\,404        & D   & AO215 & 2008-08-21  & 01:09:27.0 & +35:43:04 & 2000 & 0119+321 & 2AC & 0.78 & 256 & 0.6  & 1420.753\\tablenotemark{a} & \\nodata & 17 \\\\ NGC\\,404        & D   & AC459 & 1996-07-16  & 01:06:39.0 & +35:28:00 & 1950 & 0116+319 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & -56.0   &  0 \\\\ \\\\ KKH\\,37         & B   & AO215 & 2007-12-15  & 06:47:45.8 & +80:07:26 & 2000 & 0410+769 & 2AD & 0.78 & 256 & 0.6  & 1421.121\\tablenotemark{a} & \\nodata & 12 \\\\ KKH\\,37         & C   & AO215 & 2008-04-11  & 06:47:45.8 & +80:07:26 & 2000 & 0410+769 & 2AC & 0.78 & 256 & 0.6  & 1421.033\\tablenotemark{a} & \\nodata & 15 \\\\ KKH\\,37         & D   & AO215 & 2008-07-12  & 06:47:45.8 & +80:07:26 & 2000 & 0410+769 & 2AC & 0.78 & 256 & 0.6  & 1421.150\\tablenotemark{a} & \\nodata & 16 \\\\ KKH\\,37         & D   & AO215 & 2008-08-11  & 06:47:45.8 & +80:07:26 & 2000 & 0410+769 & 2AC & 0.78 & 256 & 0.6  & 1421.161\\tablenotemark{a} & \\nodata & 16 \\\\ \\\\ UGC\\,4483       & B   & AO215 & 2007-12-22  & 08:37:03.0 & +69:46:31 & 2000 & 0834+555 & 2AD & 1.56 & 256 &  1.3 & 1419.698\\tablenotemark{a} & \\nodata & 12 \\\\ UGC\\,4483       & B   & AZ090 & 1997-04-01  & 08:32:06.0 & +69:58:00 & 1950 & 0831+557 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & 180.0   &  0 \\\\ UGC\\,4483       & B   & AZ090 & 1997-04-11  & 08:32:06.0 & +69:58:00 & 1950 & 0831+557 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & 180.0   &  0 \\\\ UGC\\,4483       & B   & AZ090 & 1997-04-12  & 08:32:06.0 & +69:58:00 & 1950 & 0831+557 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & 180.0   &  0 \\\\ UGC\\,4483       & C   & AZ090 & 1997-06-28  & 08:32:06.0 & +69:58:00 & 1950 & 0831+557 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & 180.0   &  0 \\\\ UGC\\,4483       & C   & AZ090 & 1997-08-14  & 08:32:06.0 & +69:58:00 & 1950 & 0831+557 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & 180.0   &  0 \\\\ UGC\\,4483       & D   & AO215 & 2008-07-10  & 08:37:03.0 & +69:46:31 & 2000 & 0834+555 & 2AC & 1.56 & 256 &  1.3 & 1419.700\\tablenotemark{a} & \\nodata & 16 \\\\ UGC\\,4483       & D   & AO215 & 2008-08-16  & 08:37:03.0 & +69:46:31 & 2000 & 0834+555 & 2AC & 1.56 & 256 &  1.3 & 1419.200\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ KK\\,77          & B   & AO215 & 2007-12-07  & 09:50:10.5 & +67:30:24 & 2000 & 1035+564 & 4   & 1.56 & 128 &  2.2 & 1419.174\\tablenotemark{a} & \\nodata & 12 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.298\\tablenotemark{a} & \\nodata &    \\\\ KK\\,77          & C   & AO215 & 2008-03-31  & 09:50:10.5 & +67:30:24 & 2000 & 1035+564 & 4   & 1.56 & 128 &  2.6 & 1419.025\\tablenotemark{a} & \\nodata & 14 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.148\\tablenotemark{a} & \\nodata &    \\\\ KK\\,77          & D   & AO215 & 2008-08-10  & 09:50:10.5 & +67:30:24 & 2000 & 1035+564 & 4   & 1.56 & 128 &  2.6 & 1419.139\\tablenotemark{a} & \\nodata & 16 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.262\\tablenotemark{a} & \\nodata &    \\\\ KK\\,77          & D   & AO215 & 2008-08-19  & 09:50:10.5 & +67:30:24 & 2000 & 1035+564 & 4   & 1.56 & 128 &  2.6 & 1419.118\\tablenotemark{a} & \\nodata & 17 \\\\ &     &       &             &            &           &      &          &     &      &     &      & 1420.242\\tablenotemark{a} & \\nodata &    \\\\ \\\\ BK3N            & B   & AO215 & 2007-12-18  & 09:53:48.5 & +68:58:08 & 2000 & 1035+564 & 2AD & 0.78 & 256 & 0.6  & 1420.644\\tablenotemark{a} & \\nodata & 12 \\\\ BK3N            & C   & AO215 & 2008-03-16  & 09:53:48.5 & +68:58:08 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1420.522\\tablenotemark{a} & \\nodata & 14 \\\\ BK3N            & D   & AO215 & 2008-07-18  & 09:53:48.5 & +68:58:08 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1420.590\\tablenotemark{a} & \\nodata & 16 \\\\ BK3N            & D   & AO215 & 2008-08-15  & 09:53:48.5 & +68:58:08 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1420.631\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ AO\\,0952+69     & B   & AO215 & 2007-12-09  & 09:57:31.0 & +69:16:60 & 2000 & 1035+564 & 2AD & 1.56 & 256 &  1.3 & 1419.996\\tablenotemark{a} & \\nodata & 12 \\\\ AO\\,0952+69     & C   & AO215 & 2008-03-31  & 09:57:31.0 & +69:16:60 & 2000 & 1035+564 & 2AC & 1.56 & 256 &  1.3 & 1419.855\\tablenotemark{a} & \\nodata & 14 \\\\ AO\\,0952+69     & D   & AO215 & 2008-08-11  & 09:57:31.0 & +69:16:60 & 2000 & 1035+564 & 2AC & 1.56 & 256 &  1.3 & 1419.967\\tablenotemark{a} & \\nodata & 16 \\\\ \\\\ Sextans\\,B      & B   & AO215 & 2007-11-16  & 10:00:00.1 & +05:19:56 & 2000 & 1024-008 & 2AD & 1.56 & 256 &  1.3 & 1419.103                  & \\nodata & 12 \\\\ Sextans\\,B      & C   & AM561 & 1997-08-02  & 09:59:59.9 & +05:19:43 & 2000 & 1008+075 & 2AD & 0.78 & 128 &  1.3 & \\nodata                   & 301.0   &  0 \\\\ Sextans\\,B      & D   & AO215 & 2008-08-03  & 10:00:00.1 & +05:19:56 & 2000 & 0943-083 & 2AC & 1.56 & 256 &  1.3 & 1418.890                  & \\nodata & 16 \\\\ \\\\ NGC\\,3109       & BnA & AO215 & 2007-10-07  & 10:03:06.9 & -26:09:34 & 2000 & 0921-263 & 2AD & 1.56 & 256 &  1.3 & 1418.570                  & \\nodata & 12 \\\\ NGC\\,3109       & BnA & AO215 & 2007-10-08  & 10:03:06.9 & -26:09:34 & 2000 & 0921-263 & 2AD & 1.56 & 256 &  1.3 & 1418.570                  & \\nodata & 12 \\\\ NGC\\,3109       & CnB & AO215 & 2008-02-26  & 10:03:06.9 & -26:09:34 & 2000 & 0921-263 & 2AC & 1.56 & 256 &  1.3 & 1418.526                  & \\nodata & 14 \\\\ NGC\\,3109       & DnC & AO215 & 2008-06-15  & 10:03:06.9 & -26:09:34 & 2000 & 0921-263 & 2AC & 1.56 & 256 &  1.3 & 1418.402                  & \\nodata & 15 \\\\ NGC\\,3109       & DnC & AO215 & 2008-07-12  & 10:03:06.9 & -26:09:34 & 2000 & 0921-263 & 2AC & 1.56 & 256 &  1.3 & 1418.470                  & \\nodata & 16 \\\\ \\\\ Antlia          & BnA & AO215 & 2007-10-06  & 10:04:04.1 & -27:19:52 & 2000 & 0921-263 & 2AD & 0.78 & 256 & 0.6  & 1418.740                  & \\nodata & 12 \\\\ Antlia          & BnA & AO215 & 2007-10-13  & 10:04:04.1 & -27:19:52 & 2000 & 0921-263 & 2AD & 0.78 & 256 & 0.6  & 1418.755                  & \\nodata & 12 \\\\ Antlia          & CnB & AA232 & 1998-11-02  & 10:01:47.5 & -27:05:15 & 1950 & 1015-314 & 2AD & 0.78 & 128 &  1.3 & \\nodata                   & 360.0   &  0 \\\\ Antlia          & CnB & AA232 & 1998-11-13  & 10:01:47.5 & -27:05:15 & 1950 & 1015-314 & 2AD & 0.78 & 128 &  1.3 & \\nodata                   & 360.0   &  0 \\\\ Antlia          & DnC & AO215 & 2008-06-15  & 10:04:04.1 & -27:19:52 & 2000 & 0921-263 & 2AC & 0.78 & 256 & 0.6  & 1418.578                  & \\nodata & 15 \\\\ Antlia          & DnC & AO215 & 2008-07-26  & 10:04:04.1 & -27:19:52 & 2000 & 0921-263 & 2AC & 0.78 & 256 & 0.6  & 1418.614                  & \\nodata & 16 \\\\ \\\\ KDG\\,63         & B  &  AO215 & 2007-11-29  & 10:05:06.4 & +66:33:32 & 2000 & 1035+564 & 2AD & 0.78 & 256 & 0.6  & 1421.093\\tablenotemark{a} & \\nodata & 12 \\\\ KDG\\,63         & C  &  AO215 & 2008-04-05  & 10:05:06.4 & +66:33:32 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1420.930\\tablenotemark{a} & \\nodata & 14 \\\\ KDG\\,63         & D  &  AO215 & 2008-08-08  & 10:05:06.4 & +66:33:32 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1421.010\\tablenotemark{a} & \\nodata & 16 \\\\ KDG\\,63         & D  &  AO215 & 2008-08-17  & 10:05:06.4 & +66:33:32 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1421.052\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ Sextans\\,A      & B  &  AO215 & 2007-11-21  & 10:11:00.8 & -04:41:34 & 2000 & 1024-008 & 2AD & 1.56 & 256 &  1.3 & 1418.981                  & \\nodata & 12 \\\\ Sextans\\,A      & C  &  AO215 & 2008-03-16  & 10:11:00.8 & -04:41:34 & 2000 & 1024-008 & 2AC & 1.56 & 256 &  1.3 & 1418.799                  & \\nodata & 14 \\\\ Sextans\\,A      & D  &  AO215 & 2008-04-12  & 10:11:00.8 & -04:41:34 & 2000 & 1024-008 & 2AC & 1.56 & 256 &  1.3 & 1418.770                  & \\nodata & 16 \\\\ Sextans\\,A      & D  &  AO215 & 2008-08-17  & 10:11:00.8 & -04:41:34 & 2000 & 0943-083 & 2AC & 1.56 & 256 &  1.3 & 1418.819                  & \\nodata & 17 \\\\ \\\\ HS\\,117         & B  &  AO215 & 2007-11-28  & 10:21:25.2 & +71:06:51 & 2000 & 1035+564 & 2AD & 0.78 & 256 & 0.6  & 1420.649\\tablenotemark{a} & \\nodata & 12 \\\\ HS\\,117         & C  &  AO215 & 2008-04-11  & 10:21:25.2 & +71:06:51 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1420.500\\tablenotemark{a} & \\nodata & 15 \\\\ HS\\,117         & D  &  AO215 & 2008-07-14  & 10:21:25.2 & +71:06:51 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1420.575\\tablenotemark{a} & \\nodata & 16 \\\\ HS\\,117         & D  &  AO215 & 2008-08-11  & 10:21:25.2 & +71:06:51 & 2000 & 1035+564 & 2AC & 0.78 & 256 & 0.6  & 1420.608\\tablenotemark{a} & \\nodata & 16 \\\\ NG\\\\ DDO\\,82         & B  &  AO215 & 2007-12-13  & 10:30:35.0 & +70:37:07 & 2000 & 1035+564 & 2AD & 1.56 & 256 &  1.3 & 1420.194\\tablenotemark{a} & \\nodata & 12 \\\\ DDO\\,82         & C  &  AO215 & 2008-04-08  & 10:30:35.0 & +70:37:07 & 2000 & 1035+564 & 2AC & 1.56 & 256 &  1.3 & 1420.061\\tablenotemark{a} & \\nodata & 14 \\\\ DDO\\,82         & D  &  AO215 & 2008-08-11  & 10:30:35.0 & +70:37:07 & 2000 & 1035+564 & 2AC & 1.56 & 256 &  1.3 & 1420.165\\tablenotemark{a} & \\nodata & 16 \\\\ DDO\\,82         & D  &  AO215 & 2008-08-16  & 10:30:35.0 & +70:37:07 & 2000 & 1035+564 & 2AC & 1.56 & 256 &  1.3 & 1420.150\\tablenotemark{a} & \\nodata & 16 \\\\ \\\\ KDG\\,73         & B  &  AO215 & 2007-11-27  & 10:52:57.1 & +69:32:58 & 2000 & 1313+675 & 2AD & 0.78 & 256 & 0.6  & 1419.928\\tablenotemark{a} & \\nodata & 12 \\\\ KDG\\,73         & C  &  AO215 & 2008-04-11  & 10:52:57.1 & +69:32:58 & 2000 & 1313+675 & 2AC & 0.78 & 256 & 0.6  & 1419.778\\tablenotemark{a} & \\nodata & 15 \\\\ KDG\\,73         & D  &  AO215 & 2008-07-14  & 10:52:57.1 & +69:32:58 & 2000 & 1313+675 & 2AC & 0.78 & 256 & 0.6  & 1419.845\\tablenotemark{a} & \\nodata & 16 \\\\ KDG\\,73         & D  &  AO215 & 2008-08-16  & 10:52:57.1 & +69:32:58 & 2000 & 1313+675 & 2AC & 0.78 & 256 & 0.6  & 1419.883\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ NGC\\,3741       & B  &  AO215 & 2007-11-06  & 11:36:06.2 & +45:17:01 & 2000 & 1146+399 & 2AD & 1.56 & 256 &  1.3 & 1419.434\\tablenotemark{a} & \\nodata & 12 \\\\ NGC\\,3741       & C  &  AO215 & 2008-05-05  & 11:36:06.2 & +45:17:01 & 2000 & 1146+399 & 2AC & 1.56 & 256 &  1.3 & 1419.218\\tablenotemark{a} & \\nodata & 15 \\\\ NGC\\,3741       & D  &  AO215 & 2008-08-04  & 11:36:06.2 & +45:17:01 & 2000 & 1146+399 & 2AC & 1.56 & 256 &  1.3 & 1419.295\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ DDO\\,99         & B  &  AO215 & 2007-12-04  & 11:50:53.0 & +38:52:49 & 2000 & 1146+399 & 2AD & 1.56 & 256 &  1.3 & 1419.363                  & \\nodata & 12 \\\\ DDO\\,99         & C  &  AO215 & 2008-04-05  & 11:50:53.0 & +38:52:49 & 2000 & 1146+399 & 2AC & 1.56 & 256 &  1.3 & 1419.180                  & \\nodata & 14 \\\\ DDO\\,99         & D  &  AO215 & 2008-08-07  & 11:50:53.0 & +38:52:49 & 2000 & 1146+399 & 2AC & 1.56 & 256 &  1.3 & 1419.180                  & \\nodata & 16 \\\\ DDO\\,99         & D  &  AO215 & 2008-08-10  & 11:50:53.0 & +38:52:49 & 2000 & 1146+399 & 2AC & 1.56 & 256 &  1.3 & 1419.204                  & \\nodata & 16 \\\\ \\\\ NGC\\,4163       & B  &  AO215 & 2007-11-23  & 12:12:09.1 & +36:10:09 & 2000 & 1227+365 & 2AD & 0.78 & 256 & 0.6  & 1419.740                  & \\nodata & 12 \\\\ NGC\\,4163       & B  &  AH927\\tablenotemark{b} & 2008-02-12  & 12:12:09.1 & +36:10:09 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.719\\tablenotemark{a} & \\nodata & 13 \\\\ NGC\\,4163       & C  &  AO215 & 2008-04-08  & 12:12:09.1 & +36:10:09 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.563                  & \\nodata & 14 \\\\ NGC\\,4163       & C  &  AH927\\tablenotemark{b} & 2008-06-01  & 12:12:09.1 & +36:10:09 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.546\\tablenotemark{a} & \\nodata & 15 \\\\ NGC\\,4163       & D  &  AO215 & 2008-08-06  & 12:12:09.1 & +36:10:09 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.540                  & \\nodata & 16 \\\\ \\\\ NGC\\,4190       & B  &  AO215 & 2007-24-24  & 12:13:44.8 & +36:38:03 & 2000 & 1227+365 & 2AD & 1.56 & 256 &  1.3 & 1419.441\\tablenotemark{a} & \\nodata & 13 \\\\ NGC\\,4190       & C  &  AO215 & 2008-03-09  & 12:13:44.8 & +36:38:03 & 2000 & 1227+365 & 2AC & 1.56 & 256 &  1.3 & 1419.323\\tablenotemark{a} & \\nodata & 14 \\\\ NGC\\,4190       & D  &  AO215 & 2008-08-11  & 12:13:44.8 & +36:38:03 & 2000 & 1227+365 & 2AC & 1.56 & 256 &  1.3 & 1419.275\\tablenotemark{a} & \\nodata & 16 \\\\ \\\\ DDO\\,113        & B  &  AO215 & 2007-12-01  & 12:14:57.9 & +36:13:08 & 2000 & 1227+365 & 2AD & 1.56 & 256 &  1.3 & 1419.179                  & \\nodata & 12 \\\\ DDO\\,113        & C  &  AO215 & 2008-04-04  & 12:14:57.9 & +36:13:08 & 2000 & 1227+365 & 2AC & 1.56 & 256 &  1.3 & 1419.007                  & \\nodata & 14 \\\\ DDO\\,113        & D  &  AO215 & 2008-08-15  & 12:14:57.9 & +36:13:08 & 2000 & 1227+365 & 2AC & 1.56 & 256 &  1.3 & 1419.040                  & \\nodata & 16 \\\\ DDO\\,113        & D  &  AO215 & 2008-08-17  & 12:14:57.9 & +36:13:08 & 2000 & 1227+365 & 2AC & 1.56 & 256 &  1.3 & 1419.020                  & \\nodata & 17 \\\\ \\\\ MCG\\,+09-20-131 & B  & AO215 & 2007-11-30  & 12:15:46.8 & +52:23:17 & 2000 & 1219+484 & 2AD & 1.56 & 256 &  1.3 & 1419.749\\tablenotemark{a} & \\nodata & 12 \\\\ MCG\\,+09-20-131 & C  & AO215 & 2008-05-05  & 12:15:46.8 & +52:23:17 & 2000 & 1219+484 & 2AC & 1.56 & 256 &  1.3 & 1419.566\\tablenotemark{a} & \\nodata & 15 \\\\ MCG\\,+09-20-131 & D  & AO215 & 2008-08-10  & 12:15:46.8 & +52:23:17 & 2000 & 1219+484 & 2AC & 1.56 & 256 &  1.3 & 1419.626\\tablenotemark{a} & \\nodata & 16 \\\\ \\\\ DDO\\,125        & B  & AO215 & 2007-11-25  & 12:27:40.9 & +43:29:44 & 2000 & 1227+365 & 2AD & 0.78 & 256 & 0.6  & 1419.588                  & \\nodata & 12 \\\\ DDO\\,125        & C  & AO215 & 2008-03-09  & 12:27:40.9 & +43:29:44 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.478                  & \\nodata & 14 \\\\ DDO\\,125        & D  & AO215 & 2008-08-06  & 12:27:40.9 & +43:29:44 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.410                  & \\nodata & 16 \\\\ DDO\\,125        & D  & AO215 & 2008-08-15  & 12:27:40.9 & +43:29:44 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.446                  & \\nodata & 17 \\\\ \\\\ UGCA\\,292       & B  & AO215 & 2007-12-03  & 12:38:40.0 & +32:46:01 & 2000 & 1227+365 & 2AD & 0.78 & 256 & 0.6  & 1419.062                  & \\nodata & 12 \\\\ UGCA\\,292       & C  & AH927\\tablenotemark{b} & 2008-02-06  & 12:38:40.0 & +32:46:01 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1419.053                  & \\nodata & 13 \\\\ UGCA\\,292       & D  & AO215 & 2008-07-11  & 12:38:40.0 & +32:46:01 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1418.865                  & \\nodata & 16 \\\\ UGCA\\,292       & D  & AO215 & 2008-08-16  & 12:38:40.0 & +32:46:01 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1418.892                  & \\nodata & 17 \\\\ UGCA\\,292       & D  & AH927\\tablenotemark{b} & 2008-07-21  & 12:38:40.0 & +32:46:01 & 2000 & 1227+365 & 2AC & 0.78 & 256 & 0.6  & 1418.874                  & \\nodata & 16 \\\\ \\\\ GR\\,8           & B  & AO215 & 2007-11-12  & 12:58:40.4 & +14:13:03 & 2000 & 1254+116 & 2AD & 0.78 & 256 & 0.6  & 1419.485                  & \\nodata & 11 \\\\ GR\\,8           & C  & AH927\\tablenotemark{b} & 2008-02-10  & 12:58:40.4 & +14:13:03 & 2000 & 1347+122 & 2AC & 0.78 & 256 & 0.6  & 1419.524                  & \\nodata & 13 \\\\ GR\\,8           & D  & AO215 & 2008-08-02  & 12:58:40.4 & +14:13:03 & 2000 & 1254+116 & 2AC & 0.78 & 256 & 0.6  & 1419.285                  & \\nodata & 16 \\\\ GR\\,8           & D  & AO215 & 2008-08-17  & 12:58:40.4 & +14:13:03 & 2000 & 1254+116 & 2AC & 0.78 & 256 & 0.6  & 1419.304                  & \\nodata & 17 \\\\ GR\\,8           & D  & AH927\\tablenotemark{b} & 2008-08-02  & 12:58:40.4 & +14:13:03 & 2000 & 1347+122 & 2AC & 0.78 & 256 & 0.6  & 1419.287                  & \\nodata & 16 \\\\ \\\\ UGC\\,8508       & B  & AO215 & 2007-12-10  & 13:30:44.4 & +54:54:36 & 2000 & 1400+621 & 2AD & 0.78 & 256 & 0.6  & 1420.190\\tablenotemark{a} & \\nodata & 12 \\\\ UGC\\,8508       & B  & AH927\\tablenotemark{b} & 2008-02-09  & 13:30:44.4 & +54:54:36 & 2000 & 1400+621 & 2AC & 0.78 & 256 & 0.6  & 1420.175\\tablenotemark{a} & \\nodata & 13 \\\\ UGC\\,8508       & C  & AO215 & 2008-03-15  & 13:30:44.4 & +54:54:36 & 2000 & 1400+621 & 2AC & 0.78 & 256 & 0.6  & 1420.102\\tablenotemark{a} & \\nodata & 14 \\\\ UGC\\,8508       & C  & AH927\\tablenotemark{b} & 2008-05-31  & 13:30:44.4 & +54:54:36 & 2000 & 1400+621 & 2AC & 0.78 & 256 & 0.6  & 1420.062\\tablenotemark{a} & \\nodata & 15 \\\\ UGC\\,8508       & D  & AO215 & 2008-07-31  & 13:30:44.4 & +54:54:36 & 2000 & 1400+621 & 2AC & 0.78 & 256 & 0.6  & 1420.060\\tablenotemark{a} & \\nodata & 16 \\\\ UGC\\,8508       & D  & AH927\\tablenotemark{b} & 2008-08-03  & 13:30:44.4 & +54:54:36 & 2000 & 1400+621 & 2AC & 0.78 & 256 & 0.6  & 1420.071\\tablenotemark{a} & \\nodata & 16 \\\\ UGC\\,8508       & D  & AO215 & 2008-08-17  & 13:30:44.4 & +54:54:36 & 2000 & 1400+621 & 2AC & 0.78 & 256 & 0.6  & 1420.084\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ DDO\\,181        & B  & AO215 & 2007-12-06  & 13:39:53.8 & +40:44:21 & 2000 & 1331+305 & 2AD & 1.56 & 256 &  1.3 & 1419.543\\tablenotemark{a} & \\nodata & 12 \\\\ DDO\\,181        & C  & AO215 & 2008-03-09  & 13:39:53.8 & +40:44:21 & 2000 & 1331+305 & 2AC & 1.56 & 256 &  1.3 & 1419.474\\tablenotemark{a} & \\nodata & 14 \\\\ DDO\\,181        & D  & AO215 & 2008-08-16  & 13:39:53.8 & +40:44:21 & 2000 & 1331+305 & 2AC & 1.56 & 256 &  1.3 & 1419.388\\tablenotemark{a} & \\nodata & 17 \\\\ DDO\\,181        & D  & AO215 & 2008-08-18  & 13:39:53.8 & +40:44:21 & 2000 & 1331+305 & 2AC & 1.56 & 256 &  1.3 & 1419.390\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ DDO\\,183        & B  & AO215 & 2007-12-08  & 13:50:50.6 & +38:01:09 & 2000 & 1331+305 & 2AD & 1.56 & 256 &  1.3 & 1419.591\\tablenotemark{a} & \\nodata & 12 \\\\ DDO\\,183        & C  & AO215 & 2008-03-15  & 13:50:50.6 & +38:01:09 & 2000 & 1331+305 & 2AC & 1.56 & 256 &  1.3 & 1419.518\\tablenotemark{a} & \\nodata & 14 \\\\ DDO\\,183        & D  & AO215 & 2008-08-11  & 13:50:50.6 & +38:01:09 & 2000 & 1331+305 & 2AC & 1.56 & 256 &  1.3 & 1419.433\\tablenotemark{a} & \\nodata & 16 \\\\ DDO\\,183        & D  & AO215 & 2008-08-11  & 13:50:50.6 & +38:01:09 & 2000 & 1331+305 & 2AC & 1.56 & 256 &  1.3 & 1419.424\\tablenotemark{a} & \\nodata & 16 \\\\ \\\\ KKH\\,86         & B  & AO215 & 2007-11-11  & 13:54:33.5 & +04:14:35 & 2000 & 1347+122 & 2AD & 0.78 & 256 & 0.6  & 1419.105                  & \\nodata & 11 \\\\ KKH\\,86         & C  & AO215 & 2008-03-28  & 13:54:33.5 & +04:14:35 & 2000 & 1347+122 & 2AC & 0.78 & 256 & 0.6  & 1419.074                  & \\nodata & 14 \\\\ KKH\\,86         & D  & AO215 & 2008-08-08  & 13:54:33.5 & +04:14:35 & 2000 & 1347+122 & 2AC & 0.78 & 256 & 0.6  & 1418.944                  & \\nodata & 16 \\\\ \\\\ UGC\\,8833       & B  & AO215 & 2007-11-18  & 13:54:48.7 & +35:50:15 & 2000 & 1331+305 & 2AD & 0.78 & 256 & 0.6  & 1419.406                  & \\nodata & 12 \\\\ UGC\\,8833       & C  & AZ121 & 2000-04-15  & 13:52:38.2 & +36:04:60 & 1950 & 1413+349 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & 225.0   &  0 \\\\ UGC\\,8833       & D  & AO215 & 2008-08-05  & 13:54:48.7 & +35:50:15 & 2000 & 1331+305 & 2AC & 0.78 & 256 & 0.6  & 1419.265                  & \\nodata & 16 \\\\ UGC\\,8833       & D  & AO215 & 2008-08-15  & 13:54:48.7 & +35:50:15 & 2000 & 1331+305 & 2AC & 0.78 & 256 & 0.6  & 1419.259                  & \\nodata & 17 \\\\ \\\\ KK\\,230         & B\\tablenotemark{c} & AO215 & 2007-11-10 & 14:07:10.5 & +35:03:37 & 2000 & 1331+305 & 2AD & 0.78 & 256 & 0.6  & 1420.171\\tablenotemark{a} & \\nodata & 12 \\\\ KK\\,230         & C  & AO215 & 2008-04-03  & 14:07:10.5 & +35:03:37 & 2000 & 1331+305 & 2AC & 0.78 & 256 & 0.6  & 1420.114\\tablenotemark{a} & \\nodata & 13 \\\\ KK\\,230         & D  & AO215 & 2008-08-15  & 14:07:10.5 & +35:03:37 & 2000 & 1331+305 & 2AC & 0.78 & 256 & 0.65 & 1420.036\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ DDO\\,187        & B  & AO215 & 2007-11-17  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AD & 1.56 & 256 &  1.3 & 1419.751\\tablenotemark{a} & \\nodata & 13 \\\\ DDO\\,187        & B  & AH927\\tablenotemark{b} & 2008-02-10  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AC & 1.56 & 256 &  1.3 & 1419.805\\tablenotemark{a} & \\nodata & 13 \\\\ DDO\\,187        & B  & AH927\\tablenotemark{b} & 2008-02-12  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AC & 1.56 & 256 &  1.3 & 1419.805\\tablenotemark{a} & \\nodata & 13 \\\\ DDO\\,187        & C  & AO215 & 2008-03-28  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AC & 1.56 & 256 &  1.3 & 1419.704\\tablenotemark{a} & \\nodata & 14 \\\\ DDO\\,187        & C  & AH927\\tablenotemark{b} & 2008-05-30  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AC & 1.56 & 256 &  1.3 & 1419.674\\tablenotemark{a} & \\nodata & 14 \\\\ DDO\\,187        & D  & AH927\\tablenotemark{b} & 2008-08-05  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AC & 1.56 & 256 &  1.3 & 1419.580\\tablenotemark{a} & \\nodata & 16 \\\\ DDO\\,187        & D  & AO215 & 2008-08-06  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AC & 1.56 & 256 &  1.3 & 1419.598\\tablenotemark{a} & \\nodata & 16 \\\\ DDO\\,187        & D  & AO215 & 2008-08-16  & 14:15:56.5 & +23:03:19 & 2000 & 1330+251 & 2AC & 1.56 & 256 &  1.3 & 1419.588\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ DDO\\,190        & B  & AO215 & 2007-12-14  & 14:24:43.4 & +44:31:33 & 2000 & 1506+375 & 2AD & 0.78 & 256 & 0.6  & 1419.776                  & \\nodata & 12 \\\\ DDO\\,190        & C  & AZ121 & 2000-04-20  & 14:22:48.8 & +44:44:60 & 1950 & 1413+349 & 2AD & 1.56 & 128 &  2.6 & \\nodata                   & 160.0   &  0 \\\\ DDO\\,190        & D  & AO215 & 2008-08-16  & 14:24:43.4 & +44:31:33 & 2000 & 1506+375 & 2AC & 0.78 & 256 & 0.6  & 1419.637                  & \\nodata & 17 \\\\ \\\\ KKR\\,25         & B  & AO215 & 2008-01-31  & 16:13:47.9 & +54:22:16 & 2000 & 1634+627 & 2AD & 0.78 & 256 & 0.6  & 1421.112\\tablenotemark{a} & \\nodata & 13 \\\\ KKR\\,25         & C  & AO215 & 2008-04-03  & 16:13:47.9 & +54:22:16 & 2000 & 1634+627 & 2AC & 0.78 & 256 & 0.6  & 1421.083\\tablenotemark{a} & \\nodata & 13 \\\\ KKR\\,25         & D  & AO215 & 2008-08-07  & 16:13:47.9 & +54:22:16 & 2000 & 1634+627 & 2AC & 0.78 & 256 & 0.6  & 1421.035\\tablenotemark{a} & \\nodata & 16 \\\\ KKR\\,25         & D  & AO215 & 2008-08-16  & 16:13:47.9 & +54:22:16 & 2000 & 1634+627 & 2AC & 0.78 & 256 & 0.6  & 1421.030\\tablenotemark{a} & \\nodata & 17 \\\\ \\\\ KKH\\,98         & B  & AO215 & 2007-12-07  & 23:45:34.0 & +38:43:04 & 2000 & 0029+349 & 2AD & 0.78 & 256 & 0.6  & 1420.954\\tablenotemark{a} & \\nodata & 12 \\\\ KKH\\,98         & C  & AO215 & 2008-03-15  & 23:45:34.0 & +38:43:04 & 2000 & 0029+349 & 2AC & 0.78 & 256 & 0.6  & 1421.016\\tablenotemark{a} & \\nodata & 14 \\\\ KKH\\,98         & D  & AO215 & 2008-07-10  & 23:45:34.0 & +38:43:04 & 2000 & 0029+349 & 2AC & 0.78 & 256 & 0.6  & 1421.150\\tablenotemark{a} & \\nodata & 16 \\\\ KKH\\,98         & D  & AO215 & 2008-08-21  & 23:45:34.0 & +38:43:04 & 2000 & 0029+349 & 2AC & 0.78 & 256 & 0.6  & 1421.135\\tablenotemark{a} & \\nodata & 17  \\enddata \\tablenotetext{a}{Potential MW Interference; offset flux calibrators} \\tablenotetext{b}{additional data from the LITTLE THINGS survey} \\tablenotetext{c}{Source is named KKR 25 but is actually KK 230}   \\end{deluxetable}    \\clearpage   \\begin{deluxetable}{llrrrcccc} \\tabletypesize{\\scriptsize} \\tablecaption{Properties of the VLA-ANGST Data Cubes. \\label{tab:dataproperties}}  \\tablehead{ \\colhead{(1)}    & \\colhead{(2)}    & \\colhead{(3)}    & \\colhead{(4)}    & \\colhead{(5)}    & \\colhead{(6)}    & \\colhead{(7)}    & \\colhead{(8)}    & \\colhead{(9)}     \\\\  \\colhead{Galaxy}            & \\colhead{Weighting}         & \\colhead{B$_{major}$}       & \\colhead{B$_{minor}$}      & \\colhead{BPA}               & \\colhead{Noise}             & \\colhead{Channel Width}     & \\colhead{N$_{pixels}$}      & \\colhead{Pixel scale}      \\\\  \\colhead{}                  & \\colhead{}                  & \\colhead{[$\\arcsec$]}       & \\colhead{[$\\arcsec$]}       & \\colhead{[$\\degr$]}         & \\colhead{[mJy\\,beam$^{-1}$]} & \\colhead{[km s$^{-1}$]}     & & \\colhead{[$\\arcsec$]} } \\startdata NGC\\,247            & Natural &  9.0 &  6.2 &  10.5 & 0.9 &  2.6 & 2048$^{2}$ & 1.0 \\\\ & Robust  &  6.5 &  4.8 &  12.7 & 0.9 &      &            &  \\\\ DDO\\,6              & Natural & 12.3 & 10.3 &  52.0 & 1.9 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.2 &  6.3 &  49.0 & 2.1 &      &           &      \\\\ NGC\\,404            & Natural & 13.7 & 12.4 & -34.7 & 0.9 &  2.6 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.1 &  6.1 & -32.6 & 0.9 &      &            &      \\\\ KKH\\,37\\tablenotemark{a}             & Natural &  9.7 &  8.1 & -86.2 & 1.6 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.5 &  5.8 & -66.9 & 1.8 &      &           &      \\\\ UGC\\,4483           & Natural & 12.2 &  9.8 &  61.3 & 0.5 &  2.6 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.6 &  5.7 &  57.1 & 0.6 &      &           &      \\\\ KK\\,77\\tablenotemark{a}                & Natural & 12.2 &  8.1 & -79.0 & 0.9 &  2.6 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.1 &  5.8 & -66.6 & 0.7 &      &           &      \\\\ BK3N                & Natural & 12.0 &  8.1 & -85.1 & 1.8 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.3 &  5.8 &  61.5 & 1.8 &      &           &      \\\\ AO\\,0952+69         & Natural & 10.1 &  8.8 &  73.5 & 1.3 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.4 &  5.9 & -71.0 & 1.2 &      &           &      \\\\ Sextans~B           & Natural & 15.0 & 14.1 &  10.5 & 0.8 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  9.5 &  7.5 &  41.6 & 1.0 &      &           &      \\\\ NGC~3109            & Natural & 10.3 &  8.8 &  22.0 & 1.6 &  1.3 & 2048$^{2}$ & 1.0 \\\\ & Robust  &  7.6 &  5.0 &   8.8 & 1.7 &      &           &      \\\\ Antlia              & Natural & 14.1 & 13.9 & -81.3 & 1.0 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  & 10.5 &  9.6 &  71.1 & 1.2 &      &           &      \\\\ KDG\\,63\\tablenotemark{a}               & Natural & 10.8 &  9.2 &  85.5 & 1.4 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.2 &  6.0 &  77.0 & 1.6 &       &           &      \\\\ Sextans\\,A          & Natural & 11.8 & 10.1 &  38.5 & 1.2 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.3 &  6.0 &  35.1 & 1.3 &      &           &      \\\\ HS\\,117\\tablenotemark{a}               & Natural & 13.2 &  8.5 & -59.6 & 1.6 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  8.6 &  6.1 & -77.8 & 1.7 &      &           &      \\\\ DDO\\,82             & Natural &  9.3 &  7.7 & -81.0 & 1.3 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  5.8 &  5.7 &  65.0 & 1.4 &      &           &      \\\\ KDG\\,73             & Natural & 10.0 &  7.6 &  84.3 & 1.6 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.9 &  5.6 &  65.2 & 1.7 &      &           &      \\\\ NGC\\,3741           & Natural &  7.6 &  6.2 &  81.1 & 1.0 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  5.5 &  4.8 &  75.4 & 1.1 &      &           &      \\\\ DDO\\,99             & Natural & 12.4 &  7.6 & -86.7 & 1.0 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.7 &  5.2 &  72.5 & 1.1 &      &           &      \\\\ NGC\\,4163           & Natural & 12.3 & 10.4 & -89.6 & 1.3 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.6 &  5.4 & -85.2 & 1.4 &      &           &      \\\\ NGC\\,4190           & Natural & 10.5 &  8.9 &  83.8 & 0.9 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.1 &  5.3 &  81.4 & 1.0 &      &           &      \\\\ DDO\\,113\\tablenotemark{a}              & Natural & 15.0 & 14.0 & -55.2 & 1.4 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  9.9 &  7.7 &  82.8 & 1.5 &      &           &      \\\\ MCG\\,+09$-$20$-$131 & Natural &  9.7 &  7.4 &  69.7 & 1.0 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.1 &  5.3 &  69.1 & 1.1 &      &           &      \\\\ DDO\\,125            & Natural & 11.5 & 10.6 & -68.2 & 1.5 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.3 &  5.4 & -80.1 & 1.5 &      &           &      \\\\ UGCA\\,292           & Natural &  9.7 &  6.9 &  69.4 & 1.5 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.0 &  5.0 &  65.2 & 1.6 &      &           &      \\\\ GR\\,8               & Natural &  7.6 &  7.3 & -55.9 & 1.5 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  5.8 &  5.4 & -28.8 & 1.6 &      &           &      \\\\ UGC\\,8508           & Natural & 13.9 & 12.1 &  83.6 & 1.3 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  8.2 &  6.4 &  86.1 & 1.5 &      &           &      \\\\ DDO\\,181            & Natural & 12.8 &  9.5 & -75.7 & 1.0 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.6 &  5.5 & -80.4 & 1.1 &      &           &      \\\\ DDO\\,183            & Natural & 12.7 & 10.9 & -76.7 & 1.1 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.6 &  6.2 &  88.2 & 1.2 &      &           &      \\\\ KKH\\,86             & Natural & 11.0 &  9.9 &  -8.2 & 1.5 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.5 &  5.8 & -21.2 & 1.7 &      &           &      \\\\ UGC\\,8833           & Natural & 16.4 & 15.4 & -87.4 & 0.6 &  2.6 & 1024$^{2}$ & 1.5 \\\\ & Robust  & 12.4 & 11.2 &  81.4 & 0.6 &      &           &      \\\\ KK\\,230             & Natural &  8.2 &  7.3 & -56.5 & 1.4 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  5.9 &  5.2 & -41.6 & 1.5 &      &           &      \\\\ DDO\\,187            & Natural & 12.2 & 10.4 & -82.4 & 0.9 &  1.3 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  7.1 &  5.7 &  88.5 & 1.0 &      &           &      \\\\ DDO\\,190            & Natural & 15.6 & 14.2 &  88.1 & 0.6 &  2.6 & 1024$^{2}$ & 1.5 \\\\ & Robust  & 10.8 &  9.9 &  84.1 & 0.6 &      &           &      \\\\ KKR\\,25\\tablenotemark{a}               & Natural &  8.5 &  5.0 &  63.8 & 0.4 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  5.5 &  4.4 &  65.0 & 0.4 &      &           &      \\\\ KKH\\,98             & Natural &  9.9 &  7.4 &  82.2 & 1.3 & 0.65 & 1024$^{2}$ & 1.5 \\\\ & Robust  &  6.2 &  5.2 &  80.4 & 1.4 &      &           &      \\\\ \\enddata \\tablenotetext{a}{Non-detection}  \\end{deluxetable}       \\clearpage  \\begin{deluxetable}{lrrccccc} \\tabletypesize{\\scriptsize} \\rotate \\tablecaption{Galaxy \\hi\\ Properties. \\label{tab:hiproperties}} \\tablehead{  \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)}\\\\  \\colhead{Galaxy} & \\colhead{$S_{\\textsc{Hi}}$} & \\colhead{$M_{\\textsc{Hi}}$} & \\colhead{$S_{\\rm HI}^{\\rm SD}$} & \\colhead{$w_{\\rm 20}$} & \\colhead{$w_{\\rm 50}$} & \\colhead{$v_{\\textrm{cen}}$} & \\colhead{Peak $N_{\\textsc{Hi}}$} \\\\  \\colhead{} & \\colhead{[Jy~km~s$^{-1}$]} & \\colhead{[$10^6 M_{\\odot}$]} & \\colhead{[Jy~km~s$^{-1}$]} & \\colhead{[\\kms]} & \\colhead{[\\kms]} & \\colhead{[\\kms]} & \\colhead{[$10^{21}$\\,cm$^{-2}$]}   } \\startdata NGC\\,247          &  382.6  &  1106.2  & 608\\tablenotemark{a}  &  201.3  &  193.9  &  163.7  &   5.4  \\\\ DDO\\,6            &   1.2  &  3.2  &  3.7\\tablenotemark{b}  &  20.9  &  13.7  &  292.5  &   0.9  \\\\ NGC\\,404        \\tablenotemark{c}  &  66.7  &  146.4  &  76\\tablenotemark{d}  &  80.5  &  63.2  &  -54.0  &   0.5  \\\\ KKH\\,37\\tablenotemark{e}   &  $<3.4$   &  $<0.8$ & 1.8\\tablenotemark{f,g} & \\nodata & \\nodata & \\nodata & $<0.08$ \\\\ UGC\\,4483         &  12.0  &  32.8  &  13.6\\tablenotemark{h}  &  51.2  &  34.3  &  153.9  &   3.2  \\\\ KK\\,77\\tablenotemark{e}   & $<4.4$ & $<2.3$ & $<5.5$\\tablenotemark{i} & \\nodata & \\nodata & \\nodata & $<0.06$ \\\\ BK3N             &   6.3  &  22.0  &  $<0.75$\\tablenotemark{j}  &  44.4  &  20.0  &  -42.5  &   0.7  \\\\ AO\\,0952+69     \\tablenotemark{k}  &  61.3  &  206.6  & $<0.59$\\tablenotemark{j} &  56.0  &  45.6  &  112.8  &  1.3  \\\\ Sextans\\,B        &  91.0  &  41.5  &  72.9\\tablenotemark{h}  &  58.1  &  40.6  &  302.2  &   2.6  \\\\ NGC\\,3109         &  720.9  &  270.1  &  1148\\tablenotemark{a}  &  127.7  &  116.0  &  405.1  &   6.6  \\\\ Antlia           &   1.4  &  0.5  &  1.7\\tablenotemark{l}  &  23.4  &  13.4  &  363.4  &   0.3  \\\\ KDG\\,63\\tablenotemark{e} & $<4.2$ & $<1.1$ & $<0.2$\\tablenotemark{m} & \\nodata & \\nodata & \\nodata & $<0.06$ \\\\ Sextans\\,A        &  138.1  &  62.1  &  169\\tablenotemark{a}  &  59.8  &  46.2  &  324.8  &   6.1  \\\\ HS\\,117\\tablenotemark{e} & $<1.7$ & $<0.6$ & \\nodata  & \\nodata & \\nodata & \\nodata & $<0.03$ \\\\ DDO\\,82           &   0.8  &  2.8  &   $<0.62$\\tablenotemark{j}  &  35.8  &  26.7  &  56.2  &   0.9  \\\\ KDG\\,73           &   0.1  &  0.5  &  1.0\\tablenotemark{h}  &   9.2  &   8.5  &  116.3  &   0.1  \\\\ NGC\\,3741         &  32.8  &  81.1  &  44.6\\tablenotemark{o}  &  85.4  &  70.6  &  229.1  &   3.4  \\\\ DDO\\,99          &  29.7  &  46.9  &  47.1\\tablenotemark{o}  &  51.6&  28.6  &  242.1  &   2.6  \\\\ NGC\\,4163        &   4.8  &  9.3  &  9.6\\tablenotemark{p}  &  33.5  &  22.7  &  161.6  &   2.1  \\\\ NGC 4190         &  15.5  &  44.8  &  23.2\\tablenotemark{p}  &  73.2  &  52.8  &  227.0  &   3.5  \\\\ DDO\\,113\\tablenotemark{e} &  $<1.6$ & $<0.4$ & 23.6\\tablenotemark{p}&  \\nodata & \\nodata & \\nodata & $<0.04$ \\\\ MCG\\,+09-20-131   &   3.1  &  11.9  &  5.2\\tablenotemark{q}  &  39.0  &  26.1  &  157.6  &   3.3  \\\\ DDO\\,125          &  18.3  &  28.7  &  21.8\\tablenotemark{h}  &  39.7  &  27.0  &  196.1  &   2.1  \\\\ UGCA\\,292         &  12.9  &  40.0  &  14.3\\tablenotemark{h}  &  37.1  &  25.2  &  308.3  &   4.2  \\\\ GR\\,8             &   5.8  &  5.9  &  7.8\\tablenotemark{h}  &  30.7  &  21.4  &  213.7  &   1.7  \\\\ UGC\\,8508         &  12.3  &  19.3  &  14.8\\tablenotemark{p}  &  62.7  &  48.1  &  59.8  &   2.9  \\\\ DDO\\,181          &  10.5  &  24.4  &  12.5\\tablenotemark{o}  &  52.1  &  40.8  &  201.4  &   1.7  \\\\ DDO\\,183          &   8.2  &  20.1  &  9.6\\tablenotemark{p}  &  42.2  &  26.4  &  191.2  &   2.2  \\\\ KKH\\,86           &   0.1  &  0.1  &  0.5\\tablenotemark{h}  &   7.7  &   6.9  &  285.5  &   0.2  \\\\ UGC\\,8833         &   5.9  &  13.1  &  6.0\\tablenotemark{h}  &  41.0  &  29.4  &  225.9  &   2.2  \\\\ KK\\,230           &   0.8  &  0.7  &  2.6\\tablenotemark{h}  &  17.4  &  11.5  &  60.6  &   0.6  \\\\ DDO\\,187          &  10.1  &  11.6  &  12.0\\tablenotemark{h}  &  46.0  &  31.8  &  152.2  &   3.2  \\\\ DDO\\,190          &  22.5  &  41.3  &  8.5\\tablenotemark{o}  &  62.3  &  45.2  &  148.8  &   3.6  \\\\ KKR\\,25\\tablenotemark{e} & $<1.0$ & $<0.1$ & 2.2\\tablenotemark{r}\\tablenotemark{g} &\\nodata & \\nodata & \\nodata & $<0.03$ \\\\ KKH\\,98           &   2.2  &  3.3  &  4.1\\tablenotemark{h} &  25.5  &  17.0  &  -137.8  &   0.8  \\\\  \\enddata \\tablenotetext{a}{\\citet{huc98}} \\tablenotetext{b}{\\citet{mey04}} \\tablenotetext{c}{NGC~404 is contaminated by foreground Milky Way \\hi\\ emission.} \\tablenotetext{d}{\\citet{baa76}} \\tablenotetext{e}{VLA-ANGST non-detection. Limits based on a width of 20\\,\\kms\\ and the optical diameter $D_{25}$.} \\tablenotetext{f}{\\citet{kar01}} \\tablenotetext{g}{might be Galactic \\hi} \\tablenotetext{h}{\\citet{huc03}} \\tablenotetext{i}{\\citet{huc00a}} \\tablenotetext{j}{\\citet{vdr98}} \\tablenotetext{k}{AO~0952+62 is contaminated by M81 tidal HI emission.} \\tablenotetext{l}{\\citet{bar01}} \\tablenotetext{m}{\\citet{sch90}} \\tablenotetext{n}{\\citet{bar04}} \\tablenotetext{o}{\\citet{spr05}} \\tablenotetext{p}{\\citet{kor04}} \\tablenotetext{q}{\\citet{pus07}} \\tablenotemark{r}{\\citet{huc00b}} \\end{deluxetable}"
    },
    "1208/1208.6294.txt": {
        "abstract": "Accurately determining the properties of stars is of prime importance for characterizing stellar populations in our Galaxy. The field of asteroseismology has been thought to be particularly successful in such an endeavor for stars in different evolutionary stages. However, to fully exploit its potential, robust methods for estimating stellar parameters are required and independent verification of the results is mandatory. With this purpose, we present a new technique to obtain stellar properties by coupling asteroseismic analysis with the InfraRed Flux Method. By using two global seismic observables and multi-band photometry, the technique allows us to obtain masses, radii, effective temperatures, bolometric fluxes, and hence distances for field stars in a self-consistent manner. We apply our method to 22 solar-like oscillators in the {\\it Kepler} short-cadence sample, that have accurate {\\it Hipparcos} parallaxes. Our distance determinations agree to better than 5\\%, while measurements of spectroscopic effective temperatures and interferometric radii also validate our results. We briefly discuss the potential of our technique for stellar population analysis and models of Galactic Chemical Evolution. ",
        "introduction": "Studying the structure and evolution of the Milky Way requires detailed knowledge of the properties of the stellar populations comprising it. In this respect, \\as is a powerful tool to determine masses and radii of single stars to a high level of precision \\citep[e.g.,][]{bm10,tk10,tm10}. The {\\it CoRoT} \\citep{ab06,em08} and \\textit{Kepler} missions \\citep{gill10,wb10,dk10} have provided data on stellar oscillations of exquisite quality for thousands of stars, encouraging us to carry out a complete stellar census of the observed populations.  From the thousands of light curves obtained by the space missions, two asteroseismic parameters can be readily extracted \\citep[e.g.,][]{sh11a,dh11,wc11}. First, the power spectrum of solar-like oscillators is modulated in frequency by a Gaussian-like envelope, where the frequency of maximum power $\\nu_{\\rm{max}}$ scales approximately with the surface gravity and effective temperature. Second, the near-regular pattern of high overtones presents a dominant frequency spacing called the large frequency separation, $\\Delta\\nu$, which scales approximately with the square root of the mean stellar density. Applying scaling relations from solar values, these two asteroseismic observables may be used to estimate stellar properties of large numbers of solar-like oscillators, where individual frequencies are not available for all targets \\citep[e.g.,][]{ds08,sb10,sh11a,dh11,vsa11a}.  To gain insight about the formation history and evolution of our Galaxy, characteristics of stellar populations distributed across it must be accurately known. The best studied sample of stars in the Milky Way is the solar neighborhood, where observations and analysis over many years have determined some of its key properties, such as orbits, kinematics, and metallicities \\citep[e.g.,][]{be93,br03,bn04,fv07,fb08,lc11}. These data comprise the basic set of constraints for any model of chemical evolution of the Galaxy.  Models of Galactic Chemical Evolution are constructed under certain assumptions regarding the physical processes involved in the evolution of our Galaxy, and then calibrated against available observations. Those that reproduce them successfully are also used to predict other properties of the Galaxy, such as abundance gradients across the disk, gas infall episodes, and star formation rates \\citep[e.g.,][]{bt80,cc97,lp98,sb09}. Thus, their predictive power for our galactic history and morphology critically depends on how well they can reproduce these observations, most of which come from the solar neighborhood sample. Of particular importance among these restrictions is the age--metallicity relation, constructed using stellar isochrones and determinations of element abundances \\citep[e.g.,][]{be93,bn04}. The existence of an age--metallicity relation in the solar neighborhood is still a subject of debate \\citep[see][and references therein]{sf01,bn04,kf12}, and accurate age determinations are of prime importance to shed some new light in this issue. However, the solar neighborhood sample used to constrain these models is only complete to distances of $\\sim$50~pc \\citep[e.g.,][]{bn04}, and accurate properties of stars further than $\\sim$100~pc are difficult to measure, yet of crucial importance \\citep[e.g.,][]{fb02,mst06,zi08}. To extend the sample used as a testbed for comparison, we need stellar parameters measured with high accuracy in different regions of the Galaxy.  Asteroseismology can help bridge this gap by providing accurate stellar properties, including distances, for field stars out to several hundred parsecs. These parameters, together with effective temperatures, metallicities, and kinematics, should make possible to study spatial gradients of stellar properties across the Galactic disk and provide insight into the formation process of our Galaxy \\citep[see][]{am12a,jc12,am12c}. Moreover, robust age determinations obtained combining this information with evolutionary models will allow construction of the age-metallicity relation of the stellar populations observed by {\\it CoRoT} and {\\it Kepler}, and so provide tests outside the solar neighborhood of Galactic Chemical Evolution models.  The first comparison between asteroseismically determined parameters and predictions from Galactic Chemical Evolution models was made by \\citet{am09} for a sample of {\\it CoRoT} red giants. \\citet{wc11} used the same technique to obtain masses and radii of \\textit{Kepler} main-sequence and subgiant stars, and found a slight but statistically significant difference between the observed and synthetic mass distributions. Continuing this line of work, \\citet{am12b} obtained distances to {\\it CoRoT} and \\textit{Kepler} red giants using bolometric corrections retrieved from the literature, while \\citet{oc12} took a similar approach for five {\\it Kepler} subgiant stars.  Considering the enormous potential of the results, it is important to verify the techniques applied in asteroseismic analysis. So far, empirical tests of the scaling relations have used heterogeneous samples and relied on evolutionary models \\citep[see][and references therein]{ds09a,am12a,tbe11,mm12}. In this paper, we present a new method to derive stellar parameters in a self-consistent manner, combining seismic determinations with the InfraRed Flux Method (IRFM). We compare our results with {\\it Hipparcos} parallaxes, high-resolution spectroscopic temperature determinations, and interferometric measurements of angular diameters. We briefly discuss the implications for models of Galactic Chemical Evolution and for age determinations of main-sequence stars. % ",
        "conclusions": "\\label{conc} Determining accurate stellar parameters is crucially important for detailed studies of individual stars, as well as for characterizing stellar populations in the Milky Way. The asteroseismic revolution produced by the {\\it CoRoT} and \\textit{Kepler} missions requires robust techniques to exploit fully the potential of the data and provide the community with the building blocks for ensemble analysis.  Using oscillation data and multi-band photometry, we have presented a new method to derive stellar parameters, combining the IRFM with asteroseismic analysis. The novelty of our approach is that it allows us to obtain radius, mass, $\\teff$, and bolometric flux for individual targets in a self-consistent manner. This naturally results in direct determinations of angular diameters and distances without resorting to parallax information, further enhancing the capabilities of our technique.  Two asteroseismic methods were applied to the available data, one based entirely on scaling relations and the other one using grids of pre-calculated models. When accurate seismic data are available, comparison of our distance results with those from {\\it Hipparcos} parallaxes shows an overall agreement better than 4\\%, regardless of the asteroseismic method employed. Furthermore, the obtained $\\teff$ values show a mean difference below 1\\% when compared to results from high-resolution spectroscopy.  We have also compared our calculated angular diameters with those measured by long-baseline interferometry and found agreement within 5\\%. This provides verification of our radii, $\\teff$ values and bolometric fluxes to an excellent level of accuracy.  Despite the encouraging results, systematics can arise from faulty determinations of reddening values and metallicities. In Section~\\ref{iter} we described how the effects of metallicity and reddening, as well as their uncertainties, were taken into account in the calculations. For our determinations to be completely self-consistent, we must be able to determine those parameters from a single set of data using the IRFM. An observational campaign is currently under way to obtain Str\\\"omgren photometry of \\textit{Kepler} stars that will provide a homogeneous set of values for $\\feh$ and extinction, as described by \\citet{lc11}.  For most of the stellar parameters included in our verifications, both asteroseismic methods produce equally good results. However, it should be kept in mind that the direct method can be significantly biased when large uncertainties in the seismic input parameters exist. Moreover, scaling relations are likely to have a different dependance on effective temperature beyond the main-sequence phase, as suggested by comparison to evolutionary calculations \\citep[see][]{tw11b}. The restrictions imposed by metallicity and by the theory of stellar evolution help to cope better with large errors in seismic data, and the use of the grid-based analysis in these cases is therefore recommended .  However, to take full advantage of the available parameters, asteroseismology must provide masses with a comparable level of accuracy. It is important to note that results on masses from the direct method for values above $\\sim$~1.5~$M_\\odot$ can deviate significantly from those obtained using the grid-based approach. In fact, differences of more than $\\sim$30\\% are not unusual in these cases. Using different grids of models, \\citet{ng11} found that the fact that the direct method does not explicitly take metallicity into account could undermine its mass determinations. A thorough comparison of different grid-based techniques with the direct method is beyond the scope of this paper and will be presented in an upcoming publication (W.~J. Chaplin et al. 2012, in preparation). Another method to obtain asteroseismic masses is via detailed modeling of targets, aiming at fitting the list of individual frequencies \\citep[e.g.][]{tm10,sm12}. This approach provides mass estimates with a high level of precision and, in principle, also with high accuracy. Regardless of the considered technique, one must keep in mind that verification of asteroseismic mass determinations in general is still needed.  Studies of the stellar populations in the {\\it CoRoT} and {\\it Kepler} fields can greatly benefit from accurate masses, radii, $\\teff$, and distances \\citep{am12a,am12c}. Combining this information with evolutionary models can lead to an age-metallicity relation, opening the possibility of testing models of Galactic Chemical Evolution in stars outside the solar neighborhood \\citep[e.g.,][]{cc97,sb09,kf12}. Applying our method to the complete short-cadence {\\it Kepler} sample reveals that we can probe stars as far as 1~kpc from our Sun, making this set of main-sequence and subgiant stars extremely interesting for population studies. Although much greater distances can be probed by analyzing oscillations in giants, the ages of these stars are mostly determined by their main-sequence lifetime \\citep[e.g.,][]{ms02,sb11}. Thus, the short-cadence sample is of key importance for helping to calibrate mass-age relationships of red giants and correctly characterize their populations.  A substantial number of the {\\it Kepler} main-sequence and subgiant targets have been observed long enough to obtain individual frequency determinations \\citep{ta12}. Detailed modeling of these stars, particularly using frequency combinations\\citep[e.g.,][]{rv03,pdm10} and modes of mixed character \\citep[e.g.,][]{dm11,ob12}, can put tighter constraints on their masses and ages, providing anchor points for ensemble studies. In fact, certain combinations of frequencies can be used to probe the remaining central hydrogen content in stars \\citep[e.g.,][]{jcd88}, the existence and size of a convective core \\citep[e.g.,][and references therein]{vsa11b}, and the position of the convective envelope and helium surface abundance \\citep[e.g.,][]{jcd91,ab94}. These techniques are currently being applied to several stars in the sample (e.g., S. Deheuvels et al. 2012, in preparation, A. Mazumdar et al. 2012, in preparation, V. Silva Aguirre et al. 2012, in preparation) and should help us obtain masses with higher accuracy and determine more robust differential ages. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1208/1208.1504_arXiv.txt": {
        "abstract": "Observations of nearby galaxies have firmly established, over a broad range of galactic environments and metallicities, that star formation occurs exclusively in the molecular phase of the interstellar medium (ISM). Theoretical models show that this association results from the correlation between chemical phase, shielding, and temperature. Interstellar gas converts from atomic to molecular only in regions that are well shielded from interstellar ultraviolet (UV) photons, and since UV photons are also the dominant source of interstellar heating, only in these shielded regions does the gas become cold enough to be subject to Jeans instability. However, while the equilibrium temperature and chemical state of interstellar gas are well-correlated, the time scale required to reach chemical equilibrium is much longer than that required to reach thermal equilibrium, and both timescales are metallicity-dependent. Here I show that the difference in time scales implies that, at metallicities below a few percent of the Solar value, well-shielded gas will reach low temperatures and proceed to star formation before the bulk of it is able to convert from atomic to molecular. As a result, at extremely low metallicities, star formation will occur in a cold atomic phase of the ISM rather than a molecular phase. I calculate the observable consequences of this result for star formation in low metallicity galaxies, and I discuss how some current numerical models for H$_2$-regulated star-formation may need to be modified. ",
        "introduction": "In present day galaxies, star formation is very well-correlated with the molecular phase of the interstellar medium (ISM) \\citep{wong02a, kennicutt07a, leroy08a, bigiel08a}. In contrast, in the inner parts of disks where there are significant molecular fractions, star formation correlates very poorly or not at all with the atomic ISM. At large galctocentric radii where the ISM becomes atomic-dominated star formation does begin to correlate with H~\\textsc{i}, but this appears to be only because H$_2$ itself becomes correlated with H~\\textsc{i}, and the H$_2$ forms stars in the same way regardless of where it is found within a galaxy \\citep{bigiel10a, schruba11a}. Strong association between star formation and H$_2$ and a lack of association with H~\\textsc{i} is also found down the lowest metallicity systems that have been measured, at roughly 20\\% of Solar \\citep{bolatto11a}. In summary, all available observational data indicates that star formation occurs only where the hydrogen in the ISM has converted to H$_2$.  Theoretical models have explained these observations as resulting from a correlation between chemistry and temperature \\citep{schaye04a, krumholz11b, glover12a}. Molecular hydrogen is not an important coolant in modern-day galaxies, and while carbon monoxide (which forms only when it is catalyzed by H$_2$ -- \\citealt{van-dishoeck86a}) is, the C~\\textsc{ii} found in H~\\textsc{i} regions is almost as effective. However, H$_2$ is an excellent proxy for the presence of cold gas because both are sensitive to destruction by UV photons, which photodissociate H$_2$ and increase the temperature through the grain photoelectric effect. As a result, both H$_2$ and low temperature gas are found only in regions of high extinction where the UV photon density is far below its mean value in the ISM, and, conversely, any region that where the photodissociation rate is high enough to convert the bulk of the ISM to H~\\textsc{i} is also likely to be warm. Since low temperatures that remove thermal pressure support are a prerequisite for collapse into stars, this correlation between temperature and chemical state in turn induces a correlation between star formation and chemical state.  However, the correlation between H$_2$ and star formation must break down at sufficiently low metallicities. Before the first stars formed in the universe, and for a short time thereafter, there were no or very few heavy elements. As a result, forming H$_2$ was extremely difficult due to a lack of dust grain surfaces to catalyze the ${\\rm H~\\textsc{i}} \\rightarrow {\\rm H}_2$ reaction. Theoretical models of star formation in such environments indicate that H$_2$ fractions remain extremely small until the density rises so high ($\\gtsim 10^9$ cm$^{-3}$) that H$_2$ can form via three-body reactions \\citep{palla83a, lepp84a, ahn07a, omukai10a}. The underlying physical basis for this result is a disconnect of timescales: the equilibrium chemical state the gas would reach after a very long time would be H$_2$-dominated, but the cooling and star formation times are short enough that the gas does not reach equilibrium before collapsing into a star.  While this result has been known for zero and extremely low metallicity systems for some time, the relationship between chemical state and star formation in intermediate metallicity regime, for which observations are at least in principle possible in the local universe, has received fairly little attention. \\citet{omukai10a} consider the chemical evolution of collapsing gas cores with metallicities from 0 to Solar, and investigate under what circumstances they can form H$_2$. However, because their calculation starts with unstable, collapsing cores, it does not address the question of in what phase of the ISM one expects to find such collapsing regions in the first place, which is the central problem for understanding the observed galactic-scale correlation between ISM chemical state and star formation. \\citet{glover12b} simulate the non-equilibrium chemical and thermal behavior of clouds with metallicities from 1\\% of Solar to Solar. They find that the bulk of the cloud material converts to H$_2$ before star formation in the high metallicity clouds but not in the lowest metallicity ones, indicating that the star formation - H$_2$ correlation should begin to break down at metallicities observable in nearby galaxies. However, given the computational cost of their simulations, they are able to explore a very limited number of cases, and it is unclear how general their results might be.  The goal of this paper is to go beyond the studies of \\citet{omukai10a} and \\citet{glover12b} by deriving general results about the correlation between chemical state and star formation over a wide range of environments and metallicities. I do not perform detailed simulations, such as those of \\citeauthor{omukai10a} and \\citeauthor{glover12b}, for every case. Instead, I rely on fairly simple models that can be integrated semi-analytically. The benefit of this approach is that it is the only way to survey a large parameter space, and thereby to answer the central questions with which I am concerned: under what conditions do we expect the correlation between star formation and H$_2$ to break down? When such a breakdown occurs, what is the governing physical mechanism that causes it? What are the resulting observational signatures? What are the implications of this breakdown for the models of star formation commonly adopted in studies of galaxy formation? In the remainder of this paper, I seek to answer these questions. ",
        "conclusions": "\\subsection{Observational Implications and Tests}  \\begin{figure} \\plotone{H2ratio} \\caption{ \\label{fig:h2ratio} Ratio of H$_2$ fraction at time $t=t_{\\rm ff}$ (red) and $t=10t_{\\rm ff}$ (blue) to the equilibrium value. For each model in a grid covering the parameter range $\\log Z' = -4 - 0$, $\\log n_0 = 0 - 3$, $\\log A_V = -2 - 1$, I compute the Bonnor-Ebert mass $M_{\\rm BE}$ H$_2$ fraction at times $t = t_{\\rm ff}$ and $t=10 t_{\\rm ff}$. For this plot, I retain only models with $M_{\\rm BE}<100$ $\\msun$ at $t=t_{\\rm ff}$, indicating that these clouds are likely to form stars. As in Figure \\ref{fig:fH2Zgrid}, pixel brightness indicates what fraction of the models at a given $Z'$ fall into that particular bin of $f_{\\rm H_2} / f_{\\rm H_{2,\\rm eq}}$, with white indicating no models and solid red or blue indicating all the models. } \\end{figure}  \\begin{figure} \\plotone{tdep} \\caption{ \\label{fig:tdep} Same as Figure \\ref{fig:h2ratio}, except that the $y$ axis now represents the H$_2$ depletion time $t_{\\rm dep-H_2} = f_{\\rm H_2} (2\\mbox{ Gyr})$. } \\end{figure}  The disconnect in timescales between H$_2$ formation and cooling has two major observable consequences, which can be used as a test of the above calculations. The first of these is a drop in H$_2$ fractions below the levels predicted by equilibrium models in star-forming clouds. Figure \\ref{fig:h2ratio} shows the ratio of the H$_2$ fraction at $t=t_{\\rm ff}$ and $t=10t_{\\rm ff}$ to the equilibrium H$_2$ fraction for star-forming clouds in the model grid. Clearly we expect equilibrium models to provide good predictions for galaxies down to metallicity $Z' \\approx 0.1$ or even somewhat less. This is consistent with observations to date, which show that chemical equilibrium models provide excellent fits to observed H$_2$ to H~\\textsc{i} ratios in the Milky Way \\citep{krumholz09a, lee12a} and even the Small Magellanic Cloud (SMC; $Z'\\approx 0.2$, \\citealt{bolatto11a}). However, the Figure indicates that at metallicities of $\\log Z' = -3$, the H$_2$ fraction in a given star-forming cloud will be at most $\\sim 10\\%$ of its equilibrium value, and could be less than 1\\% of that value.  Second, the onset of star formation before the gas has time to fully transform to H$_2$ in low metallicity galaxies should manifest as a reduction in the H$_2$ depletion time $t_{\\rm dep-H_2}$, defined as the ratio of the H$_2$ mass to the star formation rate. In Solar metallicity, non-starbursting local galaxies, $t_{\\rm dep-H_2} \\approx 2$ Gyr \\citep{bigiel08a}, although lower values are possible in starbursts. This value should be lower in low metallicity galaxies by a factor of the mean H$_2$ fraction in cold, star-forming clouds, since these clouds will only partially convert to H$_2$ before forming stars and being destroyed by feedback. Figure \\ref{fig:tdep} illustrates this effect for clouds that live 1 and 10 free-fall times. Note that \\citet{glover12b} qualitatively suggested the existence of this effect, and Figure \\ref{fig:tdep} represents a quantitative extension of this prediction.  Observational tests of these predictions are complicated by the fact that H$_2$ is extremely difficult to observe at low metallicities, because CO, the traditional H$_2$ proxy, ceases to track H$_2$ at metallicities below a few tenth of Solar \\citep{krumholz11b, bolatto11a, leroy11a, shetty11b, narayanan12a, feldmann12a}. Thus direct observational tests will require the detection of H$_2$ by other means, such as dust or C~\\textsc{ii} emission that is not associated with observed H~\\textsc{i}. While observationally challenging, surveys of this sort have already been completed in the closest galaxies like the SMC \\citep{bolatto11a}, and with the observational power provided by the Atacama Large Millimeter Telescope (ALMA) should begin to be possible in even lower metallicity nearby galaxies. In particular, 850 $\\mu$m observations are an excellent probe of dust and thus all gas including H$_2$, because at 850 $\\mu$m dust is generally optically thin, the emission is not very sensitive to dust temperature, and ALMA can achieve both high spatial resolution and excellent sensitivity. Prime targets for such a campaign include IZw18, SBS 0335-052 (both $Z' \\approx 0.02$, and probably even lower dust metallicities, \\citealt{izotov99a, herrera-camus12a}), and Leo T ($Z'\\ltsim 0.01$, \\citealt{simon07a}). The ALMA observations will have to be coupled with high resolution, high sensitivity H~\\textsc{i} maps to measure the atomic content. Fortunately, sub-kpc resolution H~\\textsc{i} maps of IZw18 \\citep{van-zee98a} and SBS0335-052 \\citep{ekta09a} are already available in the literature.   \\subsection{Implications for Simulations and Semi-Analytic Models}  These results have important theoretical implications as well. Many galaxy simulation models allow star formation only in regions where the gas has converted to H$_2$; some of these models include non-equilibrium chemistry for H$_2$ formation and destruction \\citep{pelupessy09a, gnedin09a, gnedin10a, christensen12a}, while others assume equilibrium \\citep{fu10a, lagos11a, kuhlen12a, krumholz12d}. The non-equilibrium models on average yield less H$_2$ and thus less star formation at low metallicity, because often gas clouds are not able to build up significant H$_2$ fractions before being destroyed by galactic shear or similar kinematic processes \\citep{krumholz11a}. However, if the relevant timescale is the cooling time and not the H$_2$ formation time, and this effect should be far less significant. As a result, star formation should in fact occur even in gas with low H$_2$ fractions, provided that the {\\it equilibrium} H$_2$ fraction is high -- it is the equilibrium H$_2$ fraction and not the instantaneous one that correlates with gas temperature and thus is a good predictor of where star formation will occur. This suggests that, ironically, models in which the H$_2$ is assumed to be in equilibrium, while they are less accurate in predicting the actual H$_2$ fraction, may in fact be more accurate that the non-equilibrium models in predicting where star formation should occur. More generally, the calculations presented here suggest that star formation thresholds in simulations should be based on the instantaneous density and extinction, which determine the temperature, and not on non-equilibrium chemical abundances."
    },
    "1208/1208.4926_arXiv.txt": {
        "abstract": "{ Modelling isolated rotating stars at any rotation rate is a challenge for the next generation of stellar models. These models will couple dynamical aspects of rotating stars, like angular momentum and chemicals transport, with classical chemical evolution, gravitational contraction or mass-loss. Such modelling needs to be achieved in two dimensions, combining the calculation of the structure of the star, its mean flows and the time-evolution of the whole.  We present here a first step in this challenging programme. It leads to the first self-consistent two-dimensional models of rotating stars in a steady state generated by the ESTER code. In these models the structure (pressure, density and temperature) and the flow fields are computed in a self-consistent way allowing the prediction of the differential rotation and the associated meridian circulation of the stars. After a presentation of the physical properties of such models and the numerical methods at work, we give the first grid of such models describing massive and intermediate-mass stars for a selection of rotation rates up to 90\\% of the breakup angular velocity. } ",
        "introduction": "\\subsection{The astrophysical context}  In the last ten years rotation has become an essential part of stellar models. It influences all stages of stellar evolution. For instance, during  the formation process, a drastic reduction of the specific angular momentum, by a factor of order $10^5$, is observed when matter passes from the initial molecular cloud to the newly-born main-sequence star. Later, during the evolution of the star on the main sequence, many hydrodynamical instabilities, along with meridian circulations, drive some mixing in the radiative zones. The effects of this mixing are actually observed on the surface of many stars, which show elements obviously synthesized in their core region. In addition, recent work by \\cite{FHT12} shows that rotation plays a role in the nucleosynthesis of s-elements in massive stars.  Effects of rotation are also thought to be important in understanding the evolution and \u201cyields\u201d of the first generation of stars. Lacking in metals, these stars were more compact than present stars and had also weaker winds. Thus, their rotation was certainly faster than that of present day stars. It is therefore crucial to master rotational effects in order to reconstruct the history of metal enrichment in galaxies and to understand how the observed stars have been influenced by the first generation.  Lastly, we may mention an example of the importance of rotation in the end of the life of stars: recent works on the gravitational collapse of massive stars \\cite[e.g.][]{Metzger_etal11} insist on the importance of the combined effect of rotation and magnetic fields to model the explosion of supernovae and the associated production of a gamma-ray burst.  The few foregoing examples are just illustrative, since rotation influences many other aspects of stellar physics like the oscillation spectrum, mass-loss etc.  \\subsection{The 1D-models}  Presently, rotation is included in spherically symmetric stellar models through the approach proposed by \\cite{zahn92}. Since these models are one-dimensio\\-nal all of the effects of rotation are averaged on spheres. Thus differential rotation is only in the radial direction (said to be shellular) and no meridian flow appears explicitly. Because only the first terms of the spherical harmonic expansion of fields are included, this modelling is valid at small rotation rates. An equally important part of the models is that they assume the existence of some small-scale turbulence in the radiative regions that efficiently transport momentum horizontally (i.e. on isobars), erasing latitudinal gradients of angular velocity.  Although limited to slow rotation and needing some {\\it ad hoc} constants (like turbulent diffusivities), this approach has the great merit of allowing the use of 1D-models, which are now very precise as far as microphysics is concerned. It thus allowed investigators to reproduce many observed features of stars: surface abundance of lithium as a function mass \\cite[e.g.][]{CT99}, the (relative) high number of red super-giants in low-metallicity galaxies \\cite[e.g.][]{MM01}, or the ratio of type~Ibc to type~II supernovae \\cite[e.g.][]{MM05}, etc.  Although these examples reflect a truly successful modelling, the understanding of the effects of rotation is still incomplete because in many circumstances rotation is fast.  Stellar conditions thus do not meet the requirements of the models and the use of Zahn's approach becomes dubious.  Clearly, we are missing a more detailed view of reality. For instance, no one knows the rotation rates that are allowed in the foregoing 1D models.  \\subsection{The history of 2D-models of rotating stars}  The solution to the deficiencies of 1D-models about rotation will come from the use of multi-dimensional models (two-dimensional at least), which include the mean-field hydrodynamics along with the centrifugal distortion of the star. Unfortunately, building such models turned out to be a very difficult challenge. The story of this modelling is illuminating and we briefly summarize it below, following the more detailed account of \\cite{R06c}.  The first step dates back to polytropic hydrostatic models of \\cite{james64}. An important step forward was made by a US group around P.~Bodenheimer and J.~Ostriker \\cite[][]{OM68,OB68,Mark68,OH68,Jackson70,BO70,Boden71,Boden_Ost73}. Their works are based on the Self-Consistent Field method which solves Poisson's equation for the gravitational potential, $\\Delta\\phi = 4\\pi G \\rho$, with the Green function, namely  \\[ \\phi(\\vx) = -G\\int \\frac{\\rho(\\vx')}{|\\vx-\\vx'|}\\,d^3\\vx' .  \\] This solution readily includes the boundary conditions at infinity for the gravitational field.  Making general stellar models was impeded by many numerical difficulties. The codes could not deal with stellar masses less than 9~M$_\\odot$ nor with very rapid rotation. An indication of the precision reached by these models is given by the virial test (see below).  \\cite{Jackson70} found it to be around 4$\\times$10$^{-3}$.  Soon after, M.~Clement took up the challenge and solved directly the Poisson equation with a finite difference discretization \\cite[][]{clem74}. Results improved as testified by a better virial test at 2$\\times$10$^{-4}$.  Another series of work aimed at the construction of 2D-models was published by the Japanese school led by Y.~Eriguchi \\cite[see][]{Eriguchi78,ES81}.  These works led to the first attempts to account for the baroclinicity of radiative zones \\cite[][]{UE94,UE95}, but the neglecting of viscosity imposed the prescription of some {\\it ad hoc} constraints. The series ended with the work of \\cite{SHEM97} who improved the microphysics but relaxed the dynamics, considering only barotropic models.  To be complete, we need to mention the work of \\cite{EM85,EM91} who attacked the problem in a different way, using a mapping of the star. The grid points are indeed mapped to the sphere through the relation  \\[ r_i(\\theta_k) = \\zeta_i R_s(\\theta_k) \\] where $\\zeta$ is a new radial variable spanning $[0,1]$. Finite differences are used for both variables $\\zeta$ and $\\theta$. The ensuing code was quite robust, being able to compute configurations quite far from the pure sphere, but the precision of calculation as monitored by the virial test remained around $5\\times10^{-4}$.  The more recent efforts on 2D modelling not related to the ESTER project have been those of \\cite{Roxburgh04}, motivated by the need of models of rapidly rotating stars for asteroseismology, and those of \\cite{JMS04,JMS05} motivated by the discovery of the extreme flattening of the Be star Achernar \\cite[e.g.][]{DKJVONA05}. These two works use barotropic models with an imposed rotation law. We note that \\cite{JMS04} used a new version of the Self-Consistent Field technique which leads to a much more robust code, not restricted to a given mass range. These models have also been used to investigate the acoustic oscillation spectrum of rapidly rotating stars by \\cite[][]{RTMJSM09,RMJSM09}. Lastly, \\cite{deupree11} worked out a series of 2D-models with uniform rotation for masses between 1.625 and 8~M$_\\odot$. ",
        "conclusions": ""
    },
    "1208/1208.1674_arXiv.txt": {
        "abstract": "Projected rotational velocities ($v\\sin i$) are presented for a sample of 350 early B-type main sequence stars in the nearby Galactic disk. The stars are located within $\\sim 1.5$  kpc from the Sun, and the great majority within 700 pc. The analysis is based on high-resolution spectra obtained with the MIKE spectrograph on the Magellan Clay 6.5-m telescope at the Las Campanas Observatory in Chile. Spectral types were estimated based on relative intensities of some key line absorption ratios and comparisons to synthetic spectra. Effective temperatures were estimated from the reddening-free $Q$ index, and projected rotational velocities were then determined via interpolation on  a published grid that correlates the synthetic full width at half maximum of the He {\\sc i} lines at $\\lambda\\lambda4026$, 4388 and 4471 \u030a\\AA\\ with $v\\sin i$.  As the sample has been selected solely on the basis of spectral types it contains an selection of B stars in the field, in clusters, and in OB associations. The $v\\sin i$ distribution obtained for the entire sample is found to be essentially flat for $v\\sin i$ values between 0 -- 150 km s$^{-1}$, with only a modest peak at low projected rotational velocities. Considering subsamples of stars, there appears to be a gradation in the $v\\sin i$ distribution with the field stars presenting a larger fraction of the slow rotators and the cluster stars distribution showing an excess of stars with $v\\sin i$ between 70 and 130 km s$^{-1}$. Furthermore, for a subsample of potential runaway stars we find that the $v\\sin i$ distribution resembles the distribution seen in denser environments,  which could suggest that these runaway stars have been subject to dynamical ejection mechanisms. ",
        "introduction": "O and B type stars, with typical values of projected rotational velocities ($v \\sin i$) around  100 km\\,s$^{-1}$ and higher,   have the largest average $v \\sin i$ values among all main-sequence stars.  Stellar rotation appears to be a fundamental parameter constraining  the formation of these massive stars and  the environments in which they are born,  as well as their subsequent  evolution. For instance, there is observational evidence that stars formed in denser environments tend to rotate faster than those formed in associations  \\citep{wolff07} and for O and B stars in the field the proportion of slow rotators seems to be even higher (see \\citealt{hg06a} for open clusters and \\citealt{daflon07} for the Cep OB2 association). In addition, rotation may modulate the formation of massive field stars. \\cite{OL11} cite this trend, together with additional empirical evidence based on the stellar clustering law, IMF, and direct observations, as evidence that significant numbers of field massive stars form {\\it in situ}, i.e., they were not born in clusters. Also, rotation might help in understanding the origin of runaway stars. $V \\sin i$ distributions of runaway stars have not been much studied in the literature. \\cite{martin06} studied the $v\\sin i$ distribution of high latitude OB runaway stars and noted the lack of slow rotators compared to a field sample. This was interpreted in that study as evidence that those runaway stars might have been ejected from OB associations.  The study of $v \\sin i$ distributions of samples of OB stars born in different environments, such as clusters, OB associations or the general Galactic field, and selected without bias concerning cluster membership, can be used to probe the interplay between star formation and  stellar rotation. In this paper we analyse such a sample; we present the spectroscopic observations  and a first characterization of a sample of 350 OB stars located within $\\sim$ 2 kpc from the Sun. The goal of this study is to define the stars in terms of their effective temperatures, along with their projected rotational velocities, with the emphasis on the $v\\sin i$ distributions from stars in different environments. These stars will be analysed in terms of their chemical composition in a future study. This paper is divided as follow: Sect. \\ref{observation} describes the observations and sample selection; Sect. \\ref{binarity} selects from the observed sample the binary or multiple stars; Sect. \\ref{spectral} discusses the derived effective temperatures and spectral classification for the sample. Finally, projected rotational velocities are derived in Sect. \\ref{velocity}. In Sect. \\ref{discussion} we discuss the $v \\sin i$ distributions obtained for the studied sample and in Sect. \\ref{conclusions} we present the conclusions. ",
        "conclusions": "\\label{conclusions}  High resolution spectroscopic observations and a first characterization of a sample of 350 OB stars has been carried out. Projected rotational velocities were obtained for 266 stars (after rejecting spectroscopic binaries/multiple systems) using measurements of FWHM of He {\\sc i} lines and interpolation in a synthetic grid from \\cite{daflon07}. The $v\\sin i$ distribution obtained for the studied sample has a modest peak at low $v\\sin i$'s ($\\sim$  0 -- 50 km s$^{-1}$) but it is overall flat for $v\\sin i$'s roughly between 0 -- 150 km s$^{-1}$; the number of stars drops for higher values of $v\\sin i$.  The $v\\sin i$ distribution of our brighter sample stars is similar to the  one obtained from a sample of field stars picked from the work of \\cite{ALG02}.  Literature results on membership were used in order to identify subsamples of stars belonging to OB associations or clusters. We compared these two groups and found that stars members  of OB associations and clusters compose two distinct populations. The cluster stars tend to have higher $v\\sin i$'s when compared to the OB association subsample which could mean that stellar rotation of a population is dictated by the density of the cloud in which it forms. Also, when the OB association and cluster populations are compared with the field sample, it is found that the latter has a larger fraction of slowest rotators, as previously shown by other works. In fact, there seems to be a gradation from cluster to OB association to field in $v \\sin i$ distribution.  The present sample has 34 stars that were identified as runaway candidates in \\cite{tetzlaff11} catalogue.  The $v\\sin i$ distribution of the runaways sample presents two peaks: one for $v \\sin i \\sim$ 0 -- 50 km s$^{-1}$ and another for $v \\sin i\\sim$ 100 -- 150 km s$^{-1}$. The K-S test run with the runaway stars, OB association, cluster and field samples indicate that the runaway $v \\sin i$ distribution is more likely to be similar with the distribution of the denser environments, which could suggest that these stars were ejected through the dynamic ejection mechanism. Also, there is a possibility that the low $v \\sin i$ peak is composed of stars that were ejected from OB associations."
    },
    "1208/1208.4759_arXiv.txt": {
        "abstract": "Following our previous work in [JCAP {\\bf 1206}, 041 (2012) \\cite{Qiu:2012ia}], in this paper, we continue our study of reconstructing $f(R)$ modified gravity models that can be connected to a single scalar field in general relativity via conformal transformation, which lead to scale-invariant power spectrum in the early universe. With $f(R)$ modified gravity, one does not need to introduce extra scalar, the nature of which are to be explained. Different from general nonminimal coupling theory, the behavior of the $f(R)$ theory has been fixed by its counterpart in Einstein frame, and thus have one to one correspondence. Numerical plots of the functional form of $f(R)$ as well as the evolution of $R$ in terms of cosmic time $t$ are also presented. ",
        "introduction": "For theories of the early universe, the right amount of perturbations must be generated so as to conform with our observations such as cosmic miscrowave background (CMB) \\cite{Larson:2010gs} and large scale structures (LSS) \\cite{Bernardeau:2001qr}. One of the well-known observed features is, the power spectrum of these perturbations, which comes from the 2-point correlation function, has to be (nearly) scale-invariant \\cite{Larson:2010gs}, which will put on non-trivial constraints on theoretical model building. Although it is well-known that a single scalar field, which drives the universe into de-Sitter like expansion (inflation \\cite{Starobinsky:1980te,Guth:1980zm,Albrecht:1982wi,Linde:1983gd,Starobinsky:1985ww}, while the scalar is called inflaton), or nonrelativistic matter-like contraction \\cite{Finelli:2001sr,Cai:2007qw} could easily generate perturbations to meet the requirement, the nature of the scalar is still unclear.  Scale-invariant power spectrum may also arise when one modify Einstein's gravity at early times. In some cases, the modified gravity theories could be connected with unmodified general relativity (GR) plus a scalar through conformal transformations \\cite{Faraoni:1998qx}, with the latter being viewed as the counterpart in Einstein frame of the former. Due to the equivalence between the two frames (Jordan and Einstein), the perturbation generated by the couple of counterparts are exactly the same. Thanks to the connection, one can thus reconstruct models of modified gravity from the known evolution of GR plus a scalar models, which can lead to inflation or matter-contraction scenarios. Recently we proposed a way of reconstructing the models with a scalar nonminimally coupled to gravity which could give rise to scale-invariant power spectrum \\cite{Qiu:2012ia}. In this paper, we will consider another case of modified gravity, namely $f(R)$ theories. Actually as we will see later, $f(R)$ theories could be one specific but nontrivial form of nonminimal coupling. In $f(R)$ theories, there is no need to introduce the unknown scalar, and the universe is driven totally by its gravitational structure. $f(R)$ theories has been used widely as alternatives of inflation, dark matter, dark energy and so on. See \\cite{Faraoni:2000gx} for comprehensive reviews.  The reconstruction of $f(R)$ gravity has been pursued by many authors, see \\cite{Nojiri:2006be}. In their approaches, most of them reconstruct $f(R)$ theory in Jordan frame itself, provided that the cosmic evolution in Jordan frame is given. Here we will reconstruct in a different way, namely from their counterpart in Einstein frame, which looks like a single scalar field in GR, via conformal transformation. This kind of reconstruction aims at connecting different evolutions of the universe driven by modified gravity in its Jordan and Einstein frames. As is shown in \\cite{Qiu:2012ia}, in Einstein frame there are only two cases which could give rise to (nearly) scale-invariant power spectrum, namely inflation and matter-contraction. Taking the Einstein frame lagrangian as: \\be\\label{einstein} {\\cal L}_E\\sim \\frac{1}{2}R_E-\\frac{1}{2}(\\partial\\varphi_E)^2-V(\\varphi_E)~,\\ee where here and after we set the unit such that $8\\pi G=M_{Pl}^{-2}=1$, and use the metric signature $(-,+,+,+)$. A simple and representative solution is the exact solution which is obtained assuming that its equation of state $w_E$ is a constant, namely: \\bea\\label{parametrize} a_E(t_E)&\\sim&(\\pm t_E)^{\\frac{2}{3(1+w_E)}}~,~H_E(t_E)=\\frac{2}{3(1+w_E)t_E}~,\\nonumber\\\\ \\varphi_E(t_E)&=&\\frac{2\\ln(\\pm M t_E)}{\\sqrt{3(1+w_E)}}~,~V(\\varphi_E)=V_0 e^{-\\sqrt{3(1+w_E)}\\varphi_E}~\\eea where $M$ is some energy scale. In this parametrization, we have set $``+\"$ for positive $t_E$ meaning an expanding phase, while $``-\"$ for negative $t_E$ denoting a contracting phase, and $V_0$ is some constant factor. In Inflation case, we have $w_E=-1+2\\epsilon_E/3$ with the slow-roll parameter $|\\epsilon_E|\\equiv|-(dH_E/dt_E)/H_E^2|\\ll1$, then Eq. (\\ref{parametrize}) can be written as: \\bea\\label{parametrizeinf} a_E(t_E)&\\sim&{t_E}^{\\frac{1}{\\epsilon_E}}~,~H_E(t_E)=\\frac{1}{\\epsilon_E t_E}~,\\nonumber\\\\ \\varphi_E(t_E)&=&\\sqrt{\\frac{2}{\\epsilon_E}}\\ln(Mt_E)~,~V(\\varphi_E)=V_0 e^{-\\sqrt{2\\epsilon_E}\\varphi_E}~,\\eea while in matter-contraction case, one has $w_E=0$, and thus Eq. becomes: \\bea\\label{parametrizeMB} a_E(t_E)&\\sim& (-t_E)^{\\frac{2}{3}}~,~H_E(t_E)=\\frac{2}{3t_E}~,\\nonumber\\\\ \\varphi_E(t_E)&=&\\frac{2}{\\sqrt{3}}\\ln(-Mt_E)~,~V(\\varphi_E)=V_0 e^{-\\sqrt{3}\\varphi_E}~.\\eea  In this short paper, we will mainly focus on the Jordan frame of the modified gravity theories in order to find which form can be conformally connected to the above two cases, while more complete study for the case of varying $w_E$ (or $\\epsilon_E$) will be left for the future.  The remaining sections are organized as following: in Sec. II, we review the main results for the general nonminimal coupling theories that was obtained in our previous paper. in Sec. III, we focus on $f(R)$ theories. Numerical plots of the functional form of $f(R)$ as well as the evolution of $R$ in terms of cosmic time $t$ are presented. Furthermore, we also discussed about the relation of the evolutions of various cosmological variables between the two frames for an arbitrary constant $\\epsilon_E$. In Sec. IV we conclude our paper. ",
        "conclusions": "In this paper we studied the reconstruction and cosmic evolutions of $f(R)$ modified gravity models, which can be transformed as inflation or matter-contraction scenarios in their Einstein frame. The equivalence of the Jordan and Einstein frames guarantee that the perturbations generated by $f(R)$ models follows the same evolution, namely can give rise to scale-invariant power spectrum required by observations, however their background evolution might be different. In our previous work \\cite{Qiu:2012ia} we have shown that there can be more than one kind of evolution in the case of general nonminimal coupling theories, but for the $f(R)$ case, there's no such degeneracy and the correspondence between the two frames must be one to one. We find that in $f(R)$ modified gravity theory, inflation in the Einstein frame can only refer to (phantom-like) inflation in the Jordan frame, while matter-contraction in the Einstein frame can only refer to contraction with a larger equation of state in the Jordan frame. We analysed the general conditions for $f(R)$ theory of getting scale-invariant power spectrum, and obtained the evolution of the universe in the Jordan frame as well as the functional form of $f(R)$. Numerical plot of $R$ w.r.t. $t_J$ and $f(R)$ w.r.t. $R$ are also presented.  In the current paper, we only focus on $f(R)$ models corresponds to models in Einstein frame with constant $\\epsilon_E$. For case where $\\epsilon_E$ is time-varying will also be interesting, and has been investigated in many places. Varying $\\epsilon_E$ can also be one of the mechanisms of generating scale-invariant power spectrum, especially in scenarios alternative to inflation, see e.g. \\cite{Khoury:2009my}. Moreover, for whole evolution process of the universe, including reheating after inflation or transfering to late-time acceleration. For these consideration, more complicated functional form of $f(R)$ models is needed. For example, for the reheating process, other field will be introduced to interact with gravity in order to produce particles effectively. This requires new conformal relations for multi-degrees of freedom other than Eq. (\\ref{relationfr}). All these interesting topics are under investigation now.  Before ending, we would like to mention that due to the equivalence of the two frames, the Big-Bang cosmological problems (horizon, flatness, etc.) will also do no harm to the reconstructed $f(R)$ models. To see this, one can look into the efolding number $\\cal N$ defined as \\cite{Khoury:2003vb} \\be {\\cal N}\\equiv\\ln\\Big(\\frac{a_iH_i}{a_iH_i}\\Big)~,\\ee which can be directly related to these problems. Usually these problems can be avoided as long as we require that ${\\cal N}\\gtrsim 70$ during inflation. From the relation (\\ref{relationfr}) we can see that the conformal Hubble parameter, ${\\cal H}\\equiv aH$, is not conformal invariant, but since in our case $\\delta_F$ is a constant, ${\\cal N}$ is a conformal invariant variable. Therefore, provided that inflation lasts for enough efolding number in one frame, one need not worry about whether it does in the other frame. We'd also like to refer the readers to \\cite{Qiu:2012ia} for more detailed arguments."
    },
    "1208/1208.4273_arXiv.txt": {
        "abstract": "It is assumed that accreting neutron stars in low-mass X-ray binaries are heated due to the compression of the existing crust by the freshly accreted matter which gives rise to a variety of nuclear reactions in the crust.  It has been shown that most of the energy is released deep in the crust by pycnonuclear reactions involving low-Z elements (the deep-crustal heating scenario). In this paper we discuss if neutron stars in so-called very-faint X-ray transients (VFXTs; those which have outburst peak 2-10 keV X-ray luminosities $< 1\\times 10^{36}$ erg s$^{-1}$) can be used to test this deep-crustal heating model. We demonstrate that such systems would indeed be very interesting objects to test the deep-crustal heating model with, but that the interpretation of the results might be challenging because of the large uncertainties in our estimates of the accretion rate history of those VFXTs, both the short-term (less than a few tens of thousands of years) and the one throughout their lifetime. The latter is particularly important because it can be so low that the neutron stars might not have accreted enough matter to become massive enough that enhanced core cooling processes become active. Therefore, they could be relatively warm compared to other systems for which such enhanced cooling processed have been inferred. However, the amount of matter can also not be too low because then the crust might not have been replaced significantly by accreted matter and thus a hybrid crust of partly accreted and partly original, albeit further compressed matter, might be present. This would inhibit the full range of pycnonuclear reactions to occur and therefore possibly decreasing the amount of heat deposited in the crust. More detailed calculations of the heating and cooling properties of such hybrid crusts have to be performed to be conclusive. Furthermore, better understanding is needed how a hybrid crust affects other properties such as the thermal conductivity. A potential interesting way to observe the effects of a hybrid crust on the heating and cooling of an accreting neutron star is to observe the crust cooling of such a neutron star after a prolonged (years to decades) accretion episode and compare the results with similar studies performed for neutron stars with a fully accreted crust.  We also show that some individual neutron-star low-mass X-ray binaries might have hybrid crusts as well as possibly many of the neutron stars in high-mass X-ray binaries. This has to be taken into account when studying the cooling properties of those systems when they are in quiescence. In addition, we show that the VFXTs are likely not the dominate transients that are associated with the brightest ($\\sim 10^{33}$ erg s$^{-1}$) low-luminosity X-ray sources in globular clusters as was previously hypothesized. ",
        "introduction": "Neutron-star low mass X-ray binaries (LMXBs) harbor neutron stars which are accreting matter from a close-by low-mass (typically $<1$ M$_\\odot$) companion star which transfer mass due to Roche-lobe overflow.  In most systems the neutron star does not continuously accrete matter. Generally, those systems (called X-ray transients) are in their quiescent state in which they do not accrete at all or only at a very low rate. Only occasionally they exhibit bright X-ray outbursts during which their X-ray luminosities increase by several orders of magnitude. Such outbursts are most probably caused by a very large increase in the mass accretion rate onto the neutron stars due to instabilities in the accretion disc \\citep[see the review by][]{2001NewAR..45..449L}. The X-ray transients can be divided in sub-groups based on their peak 2--10 keV X-ray luminosities in outburst. In particular, a special sub-group has been called the very-faint X-ray transients (VFXTs) that have peak luminosities between $1\\times 10^{34}$ and $1\\times 10^{36}$ erg s$^{-1}$ \\citep[see][]{2006A&A...449.1117W}, whereas the brighter transients have peak luminosities of $10^{36-39}$ erg s$^{-1}$.  In quiescence, neutron-star transients can still be detected using sensitive X-ray satellites, and it has been found that in many systems a soft, most likely thermal, component is present with a typical black-body temperature of 0.1--0.3 keV \\citep[see, e.g.,][and references to those papers]{1987A&A...182...47V,1996PASJ...48..257A,1998A&ARv...8..279C,1999ApJ...514..945R}. In addition, for many systems an additional spectral component above 2 keV has also been detected \\citep[the non-thermal power-law component; see, e.g.,][]{1996PASJ...48..257A,1999ApJ...514..945R}, which can even dominate the 0.5--10 keV X-ray flux in some systems \\citep[e.g.,][]{2002ApJ...575L..15C,2004MNRAS.354..666J,2005ApJ...618..883W,2012arXiv1204.6059D}. The origin of the non-thermal component is not clear \\citep[see, e.g., the discussions in][]{1998A&ARv...8..279C,2012arXiv1204.6059D} but it is generally assumed that the soft component is the thermal emission from the neutron star surface, either due to very-low level residual accretion onto the surface or due to the cooling of the neutron star that has been heated by the matter accreted during the outbursts.  During the accretion phases, matter accumulates on the surface of the neutron star. This matter compresses the underlying layers of the neutron star crust.  If the accretion continues long enough the original catalyzed crust can be completely replaced by a new crust made of accreted matter (\\citealt{1979PThPh..62..957S}, \\citealt{1990A&A...229..117H}).  The original crust is pushed down into the neutron star until it fuses together with the core. The composition of the accreted crust should be quite different from the original, catalyzed crust, i.e., richer in low-Z elements \\citep{1990A&A...229..117H}. It has been postulated that when the accreted matter sinks into the crust due to the compression induced by freshly accreted material onto the star, a chain of non-equilibrium reactions occurs in the crust that generates heat \\citep[electron captures, neutron drips and pycnonuclear reactions;][]{1990A&A...227..431H,2003A&A...404L..33H,2008A&A...480..459H,2007ApJ...662.1188G}. Most of the heat is released deep in the crust (at densities $>10^{12}$ g cm$^{-1}$) due to pycnonuclear reactions involving low-Z elements. This heat is conducted inwards, heating the core, and outwards, where it is emitted as thermal emission from the surface. This model has been called the ``deep crustal heating model'' \\citep{1998ApJ...504L..95B}. This model has been tested by comparing the observed thermal emission of quiescent neutron stars with predictions based on estimated of their time-averaged accetion rates (see section~\\ref{Sec:crustheating_data}). Another exciting possibility is to study the thermal relaxation of accretion-heated neutron star crusts after the end of accretion outbursts (see also section \\ref{forward}).  In Section~\\ref{Sec:crustheating} we briefly describe the deep crustal heating model and compare the model with the available data. We also calculate the time-scale on which the core reacts to changes in the long-term averaged accretion rate. In Section~\\ref{Sec:qLx} we calculate, in the frame work of the deep-crustal heating model, the expected quiescent luminosity of neutron-star VFXTs, in order to use those systems to test the deep-crustal heating model. We argue that it might be possible that during their life those systems might not have accreted enough matter to have fully replaced their original neutron-star crust with an accreted one which could significantly inhibit the pycnonuclear heating reactions.  In Section~\\ref{sec:discussion} we discuss how the VFXTs can still be used to test the model and also discuss potential other sources which might harbor neutron stars with only partly accreted crusts. ",
        "conclusions": "\\label{sec:discussion}   We have estimated the quiescent thermal luminosity of neutron-star VFXTs in order to determine if they can be used to test the deep-crustal heating model in an hardly explored \\averagedmdot~regime. Unfortunately a conclusive answer cannot be give due to the large uncertainties in our knowledge of the accretion rate history of VFXTs. The \\averagedmdot~of the source during the last several thousands to tens of thousand years determines how much heat has be deposited in the neutron star over that period and therefore the thermal state of the star. However the long-term history over the lifetime of the binary determines the amount of matter accreted and therefore if enough matter has been accreted to trigger enhanced neutrino emission processes in the core and if enough matter is accreted to allow the activation of all pycnonuclear heating reactions in the inner crust.  This last point arises because it is well possible that the amount of matter which primordial VFXTs have accreted during their lifetime is not enough to fully replace the original crust, leaving a crust which is partly replaced by accreted matter and partly still contains the original, albeit further compressed material. It is unclear how such a hybrid crust would react to the accretion of matter and how this would effect the thermal state of the neutron star. Likely less heat is produced because not all pycnonuclear reactions can occur, but it is not clear if other properties of a hybrid crust are also significantly different compared to a fully accreted crusts, such as the thermal conductivity. Besides obtaining more observational data to constrain the models, detailed theoretical calculations have to be performed to investigate the heating and cooling in neutron stars which have a hybrid crust. In particular it is important to investigate different crustal compositions with a variety of amount of matter accreted \\cite[e.g., an update study of the one performed by][]{1979PThPh..62..957S}. This problem might not only be interesting for VFXTs, but also to other types of neutron stars because VFXTs might not be the only sources which harbor neutron stars with hybrid crusts (see Section~\\ref{additional}). Furthermore, VFXTs might be important to understand low-luminosity X-ray sources in globular clusters (Section~\\ref{globulars}). In addition, they might form an interesting group of sources to try to study cooling of the neutron star crust (Section~\\ref{forward}) after it has been heated during outbursts.  \\subsection{Additional potential sources without fully accreted crusts \\label{additional}}  Despite that it is generally accepted that most neutron-star LMXBs are rather old systems with ages of $10^{8-9}$ year, there are individual sources which likely are much younger.  One relatively young system might be the recently discovered transiently accreting  11 Hz X-ray pulsar IGR J17480--2446 in the globular cluster Terzan 5 \\citep[also called Terzan 5 X-2;][]{2010ATel.2929....1S,2011A&A...526L...3P}. This system is an unusual LMXB because it was expected that the neutron stars in LMXBs should have spin periods $<10$ millisecond because they are spun up by the accretion of matter \\citep[see review by][]{1991PhR...203....1B}. The slow spin period of IGR J17480--2446 is enigmatic and it has been hypothesized that this is due to the fact that the system has so far only spend a relatively brief time in the Roche-lobe overflow phase \\citep[$10^7$ to $10^8$ years;][]{2012ApJ...752...33P}. If true, this system might be an example of systems which do not have a fully replaced neutron star crust.  The mass accretion rate of this source during outburst has been estimated to be $3 \\times 10^{-9}$ M$_\\odot$ yr$^{-1}$ \\citep{2011MNRAS.412L..68D}. The duty cycle of this system is poorly constrained but if we assume again values of 1\\%-10\\% we obtain a time-averaged accretion rate of $3 \\times 10^{-11}$ - $3\\times 10^{-10} $ M$_\\odot$ yr$^{-1}$. Combined with the expected age of the accretion phase this results in a mass accreted on the neutron star of $3\\times 10^{-4}$ to $3\\times 10^{-2}$ M$_\\odot$. Although the maximum amount of matter accreted would indicate that the full crust is replaced, it is also quite possible that the neutron star in this system has a hybrid crust as well. \\cite{2011MNRAS.412L..68D} found the quiescent counterpart for this source to be rather cold, significantly colder than expected using standard heating and cooling theory. They suggested that in its neutron star enhanced core cooling processes might be active although, as also shown above, probably not enough matter has accreted on the star for the star to have become massive enough to allow such processes to occur in the core. Alternatively, they suggested that the duty cycle might be extremely low, of the order 0.1\\%.  Although not impossible, this duty cycle seems very low (and possibly improbable in the disk instability model) and therefore we suggest an another possible reason why the source is so faint in quiescence\\footnote{We note that the thermal quiescent luminosity of the source is still well within the range observed from other quiescent neutron-star LMXBs which would suggest that the source is not special. This could indicate that the same physical processes are at work in this source as well as in the other sources. This would also satisfy the principle of Occam's razor, by not having to have to postulate several mechanisms why certain quiescent LMXBs are colder than expected by the standard theory.}: due to the presence of a hybrid crust, not all the heating reactions can occur in the crust and therefore less heat has been deposited in the neutron star to heat it up to the expected temperature as inferred from its \\averagedmdot. This conclusion still holds when also taking into account that before the Roche-lob overflow phase a wind-accretion phase occurred. \\cite{2012ApJ...752...33P} estimated that the mass accretion rate in that phase would at most be $10^{-13}- 10^{-12}$ M$_\\odot$ yr$^{-1}$. This phase could have lasted $10^{7-8}$ year and thus at most $10^{-6} - 10^{-4}$ M$_\\odot$ could have been accreted.   If this is the correct explanation for why the neutron star in IGR J17480--2446 is colder than expected, one has to wonder if a similar argument might also hold for other systems which have found to be too cold. For example, the neutron star in SAX J1808.4--3658 seems to be extremely cold \\citep[][]{2002ApJ...575L..15C,2007ApJ...660.1424H,2009ApJ...691.1035H}. Its \\averagedmdot~has been estimated to be $\\sim 10^{-11}$ M$_\\odot$ yr$^{-1}$ \\citep{2007ApJ...660.1424H} and if it has lived shorter than $10^{8}$ years the neutron star should have a hybrid crust. However, this system is an accreting millisecond pulsar with a spin period of 401 Hz \\citep{1998Natur.394..344W}. This means that a significant amount of matter has to have been accreted by the neutron star to spin it up to this spin frequency. Typically, the calculations show that at up to 0.1 M$_\\odot$ \\citep{1987IAUS..125..393V} is needed to accomplish this \\citep[see review by][]{1991PhR...203....1B}. Therefore, in SAX J1808.4--3658 the neutron-star crust will have been fully replaced, which strongly indicates that in the past the accretion rate of this system was considerably larger than its current inferred \\averagedmdot. Another source which might be relatively young is Circinus X-1. The age of this system is not known, but it has been suggested to be rather young \\citep[of the order of $<10^{4-5}$ years; see the discussion in][]{2004MNRAS.348..458C}. Despite that it can accrete on occasions at very high accretions rates (up to $>10^{-8}$ M$_\\odot$ yr$^{-1}$), this age (if confirmed) is sufficiently low that very likely not the complete crust has been replaced. If the source would go fully quiescence, it would be very interesting to determine the quiescent luminosity of the neutron star in this system.  \\subsubsection{Neutron stars in high-mass X-ray binaries}  In high-mass X-ray binaries (HMXBs) the neutron star accretes either from the strong stellar wind of the companion (e.g., a supergiant star) or from the decretion disk of a B type star, which is typically observed to be of type B0-B2 in Be/X-ray transients \\citep[see review by][]{2011Ap&SS.332....1R}. Such early-type B stars only live between 10 to 30 million years. Typically in Be/X-ray transients the sources have outbursts with X-ray luminosities of $10^{36-37}$ erg s$^{-1}$ (corresponding to an outburst accretion rate of $10^{-10}$ to $10^{-9}$ M$_\\odot$ yr$^{-1}$) when the neutron star moves through the decretion disk at periastron passage.  For some sources this occurs once every orbital period (resulting in periodic outbursts; called type-I outburst) but other sources are only occasionally in outburst. Therefore, it is unclear what fraction of the time the neutron star is actually accreting, but when assuming again a duty cycle of 1\\%-10\\%, this would result in a \\averagedmdot~of $10^{-12}$ to $10^{-10}$ M$_\\odot$ yr$^{-1}$ and a total amount of mass accretion throughout the life time of the system (assuming the system was a Be/X-ray transient for the whole life of the B star which might be a significantly overestimation of the duration of this phase) of $10^{-5}$ to $3 \\times 10^{-3}$ M$_\\odot$. Thus, it is quite possible that also the neutron stars in some Be/X-ray transients have a hybrid crust. We note that some systems also exhibited so-called type-II outbursts which are much brighter (peak luminosities of $10^{38}$ erg s$^{-1}$) which can last for weeks to months but they are very infrequent and not all systems exhibit them. Therefore, we do not expect that those type of outbursts will affect our main conclusion significantly.  Another class of HMXBs transients are the supergiant fast X-ray transients \\citep[or SFXTs; see, e.g.,][]{2011AdSpR..48...88S} in which the neutron star transiently accretes from the variable dense wind of a supergiant star. However, also very likely in those systems the neutron star has only a partly replaced crust because the supergiants only live very short and despite that the outbursts of those systems can be very bright ($10^{38}$ erg s$^{-1}$) they are very brief, very infrequent, and most of the time the neutron star is only accreting at much lower rates or not at all. Although, if before the supergiant phase the neutron star was also already accreting significantly from the companion star \\citep[e.g., during an earlier Be phase;][]{2011MNRAS.415.3349L}, then more of the original crust is replaced.  For the neutron stars in HMXB transients~\\averagedmdot~is typically higher than inferred for the VFXTs (they are typically more in the range observed for the ordinary LMXB transients). Therefore, it is expected that if standard heating and cooling occurs in those systems \\citep[as suggested by][]{1998ApJ...504L..95B}, that their thermal emission should be readily detectable in quiescent. Enhanced core cooling is not expected because they should be relatively light weight neutron stars since little matter has been accreted \\citep[although it might be possible that some systems are born with massive neutron stars; see, e.g.,][]{2001A&A...377..925B}. In contrast, the heating might be affected by what type of crust is present (i.e., fully accreted or hybrid crust) and HXMB transients might be very good candidates to investigate the effect of hybrid crust on the thermal properties of the neutron star. However, the situation for those sources might be complicated by the much stronger magnetic field in those systems ($10^{12-13}$ Gauss) compared to those of the neutron stars in LMXBs ($10^{8-9}$ Gauss). It is unclear how strong the effects of these stronger magnetic fields are on the heating and cooling of the neutron stars and other related properties \\citep[e.g., the thermal conductivity which is severally affects by super strong magnetic fields of $>10^{13}$ Gauss and therefore likely also by slightly lower fields;][]{1999A&A...346..345P,2008A&A...486..255A}. More detailed theoretical calculations have to be performed to determine the effect of the magnetic field, in combination with the exact composition and structure of the (possible hybrid) crust.  Observing HMXB transients in their quiescent state could be very useful in this aspect. However, the number of neutron-star Be/X-ray transients so far studied in quiescence is rather limited \\citep[for a source list see][]{2007ApJ...658..514R,2011ApJ...728...86T}. So far, the obtained picture is complex. Some systems (like, e.g., EXO 2030+375) always remain rather bright in-between outbursts ($>10^{35}$ erg s$^{-1}$; basically they never transit to quiescence). However, the majority of systems have quiescent luminosities between $10^{32}$ and $10^{34}$ erg s$^{-1}$ \\citep[][]{2007ApJ...658..514R,2011ApJ...728...86T}. Spectral analysis demonstrates that some systems are still very hard in quiescence with power-law indices near 1 or even lower \\citep[similar to what often is seen in outburst;][e.g.,]{2007ApJ...658..514R}, while others are softer with indices even up to 2.6 \\citep[e.g.,][]{2002ApJ...580..389C}.  Although the quiescent data are usually not of very high quality, several sources do not show pulsations in quiescence which might indicate that indeed the accretion down to the surface has halted in those systems \\citep[][]{2002ApJ...580..389C,2005ApJ...622.1024W}. However, in a few other systems pulsations could still be detected in quiescence demonstrating that in those systems either some of the matter still reaches the neutron-star surfaces or the pulsations are in some way caused by the interaction between the magnetic field (which is rotating with the neutron star) and the accretion of matter down to the magnetosphere \\citep[][]{2007ApJ...658..514R,2005A&A...431..667M}.  For some systems it has been suggested \\citep[][]{2002ApJ...580..389C} that indeed the emission we observe is due to the cooling of the neutron star and not due to some sort of accretion process, however, the evidence is not conclusive due to the statistical quality of the data. Furthermore, the possibility that they might harbor a neutron star with a hybrid crust was not considered. \\cite{2011ApJ...728...86T} discussed possible reason why the candidate Be/X-ray transient IGR J01363+6610 could not be detected with {\\it Chandra} in its quiescent state (e.g., the system containing a black hole instead of a neutron star), but in light of the above discussion we suggest that the possibility should be considered that this source might still harbor a neutron star but one with a hybrid crust which inhibits significant heating of the neutron star.  The situation for SFXTs is similar to that of the Be/X-ray transients with only a handful of SFXTs studied in quiescence. Also those systems show a variety in quiescent behavior \\citep[see, e.g.,][]{2005A&A...441L...1I,2010A&A...519A...6B,2012arXiv1207.3719B}. A systematic and homogenous study of many more HMXB transients (both SFXTs and Be/X-ray transients) in quiescence is needed to understand fully how they can be used to study the deep-crustal heating model. A survey (using {\\it Chandra}) of 16 confirmed neutron-star Be/X-ray transients in their quiescent state has recently been accepted (PI: Wijnands) which will give more insight into this issue.   \\subsection{VFXTs in globular clusters \\label{globulars}}  Many faint X-ray sources have been found in the Galactic globular clusters, and a large number are likely associated with neutron-star X-ray transients \\citep[see, e.g.,][]{1984MNRAS.210..899V}. But the lack of a significant number of outbursts from those sources has lead to suggestions that maybe those sources are associated with VFXTs whose outbursts where missed by the all-sky monitors \\citep[][]{2008AIPC.1010..382W}. As estimated in section~\\ref{qLxestimate}, the quiescent X-ray luminosity of VFXTs in the standard deep-crustal heating model would be in the range $10^{31-33}$ erg s$^{-1}$ and indeed, if the standard heating and cooling processes occur, a large fraction of the candidate quiescent LMXBs could be associated with VFXTs. However, as also explained in section~\\ref{enhanced} and~\\ref{hybrid} the quiescent luminosity of VFXTs could be significantly lower than expected in the standard model and therefore it is unclear if this conclusion still holds. Moreover, even in the standard model it is unlikely that the VFXTs are associated with the candidate quiescent LMXBs in globular clusters.  To demonstrate this, we rewrite the time-averaged accretion rate into  \\begin{equation} \\langle \\dot{M} \\rangle = {\\langle \\dot{M}_o \\rangle t_o + \\langle \\dot{M}_q \\rangle t_q \\over t_o + t_q} \\approx   \\langle \\dot{M}_o \\rangle {t_o  \\over t_o + t_q} =  \\langle \\dot{M}_o \\rangle~ {\\rm DC} \\label{DC} \\end{equation}  \\noindent with $\\langle \\dot{M}_o \\rangle$ the time-averaged accretion rate in outburst and $\\langle \\dot{M}_q \\rangle$ the time-averaged accretion rate in quiescence. Equation \\ref{DC} assumes that $\\langle \\dot{M}_o \\rangle t_o \\gg \\langle \\dot{M}_q \\rangle t_q$, which is usually true but might not if $t_q \\gg t_o$, thus for systems with a extremely low duty cycle. Using equations~\\ref{Eq:Heat} and \\ref{DC}, and assuming $L_q = \\langle H \\rangle$ (thus standard, slow cooling) one obtains  \\begin{eqnarray} DC = { L_q \\over  10^{33} {\\rm ~erg~ s}^{-1}} { 10^{-11}~ M_\\odot ~{\\rm yr}^{-1} \\over\\langle \\dot{M}_o \\rangle } {  1.5 ~{\\rm MeV}\\over Q_{nuc}}\\\\ >{ L_q \\over 3 \\times 10^{34} {\\rm ~erg~ s}^{-1} } {  1.5 ~{\\rm MeV}\\over Q_{nuc}}\\label{eqDC} \\end{eqnarray}  \\noindent which assumes $ \\langle \\dot{M}_o \\rangle <3 \\times 10^{-10}$ $M_\\odot$ yr$^{-1}$ (the limit set by the all-sky monitors).    Typically quiescent LMXB candidates in globular clusters have bolometric X-ray luminosities of $10^{32-33}$ erg s$^{-1}$ with the brightest being $\\sim 3.6 \\times 10^{33}$ erg s$^{-1}$ \\citep[][about a third of the sources listed in that paper have an X-ray luminosity $>1 \\times 10^{33}$ erg s$^{-1}$]{2003ApJ...598..501H}. This results (using equation~\\ref{eqDC}) in limits on the duty cycle of DC $>$ 0.0025 - 0.005 for $L_q = 1 \\times 10^{32}$ erg s$^{-1}$, DC $>$ 0.025 - 0.05 for $L_q = 1 \\times 10^{33}$ erg s$^{-1}$ and DC $>$ 0.1 - 0.2 for $L_q = 3.6 \\times 10^{33}$ erg s$^{-1}$. The range in DC is due to the fact that we have assumed that $Q_{nuc} $ must be between 1 and 2. Putting reliable observational constraints on the duty cycle of possible VFXTs in globular clusters is difficult. However, currently there are about 30 quiescent LMXB candidates identified for which no outbursts\\footnote{This excludes the two sources (IGR J17480--2446, Swift J174805.3--244637) in Terzan 5 which were previously identified as candidate quiescent LMXBs \\citep[][]{2006ApJ...651.1098H} but which have now been shown to be associated with bright transients. However, the final conclusions does not depend very sensitive on how many quiescent LMXBs are currently known but have not been associated yet with outbursts.}  have been seen yet (see \\cite{2003ApJ...598..501H} for a list, several additional sources have found since that publication; e.g., \\cite{2007ApJ...657..286L,2009MNRAS.392..665G,2011ApJ...738..129G,2012arXiv1208.1426M}). During the {\\it Chandra} and {\\it XMM-Newton} observations which detected those sources, they were not in outburst. Therefore, assuming that all sources are from the same population of transient type, the duty cycle of those systems is $<1/30 = 0.033$ \\cite[see also][]{2010AIPC.1314..135H}. We note that a number of clusters have multiple {\\it Chandra} and {\\it XMM-Newton} observations during which those sources were not seen in outbursts. However, those additional observations do not constrain the duty cycle further because it is quite possible that the quiescent duration is significantly longer than the sampling time scale meaning that the observations probe the same quiescent period and therefore are not statistically independent (basically many outburst-quiescent cycles must have passed for the sampling, if performed with random time-intervals between the observations, to become independent). We note that the uncertainties on the duty cycle inferred from observations is likely to be large, but the duty cycle has to be $>$10\\% to explain the brightest sources (in the standard deep-crustal heating model) which seems unlikely. Furthermore, similar low duty cycles were inferred for VFXTs near the Galactic center \\citep{2010A&A...524A..69D}, so it is quite possible that they indeed have such low duty cycles.   The conclusion from the above exercise is that VFXTs can still be associated with some of the quiescent LMXB candidates, but only with the faintest sub-set of the group and likely not with the brightest objects ($> 10^{33}$ erg s$^{-1}$; especially not those of a few times $10^{33}$ erg s$^{-1}$; note that this assumed standard heating and cooling). This is supported by the fact that the three VFXTs currently known in globular clusters \\citep[M15 X-3, NGC 6440 X-2, and NGC 6388;][]{2009ApJ...692..584H,2010ApJ...714..894H,2011A&A...535L...1B} all have quiescent luminosities (well) below $10^{32}$ erg s$^{-1}$ \\citep[only M15 X-3 has been detected;][]{2009ApJ...692..584H}. It seems that such faint transients are indeed present in globular clusters, but that their quiescent luminosities is really low. It is worth noting that those are the only systems for which useful constraints have been set on the quiescent properties of neutron-star VFXTs. The Galactic disk sources are usually too much absorbed (see also Section~\\ref{forward}) for any useful constraints indicating that globular clusters are prime targets to study quiescent VFXTs. Sensitive monitoring programs \\citep[see][]{wijnandsatel,altamirano2012} are needed to detect the very-faint outbursts of VFXTs in globular clusters and combined with rapid follow-up observations using {\\it Chandra} \\citep[][]{altamirano2012} to determine the exact position, the quiescent counterpart can then be studied in archival {\\it Chandra} data or in newly proposed {\\it Chandra} observations.  It might be possible that we have underestimated the amount of energy released per accreted nucleon \\citep[see, e.g.,][]{2012PhRvC..85e5804S} and therefore the systems should be more luminous per accreted nucleon than assumed \\citep[e.g., evidence has been reported that in some systems extra heat in shallower layers in the crust must be liberated to explain the quiescent properties of those systems; ][]{2009ApJ...698.1020B,2011MNRAS.418L.152D}. However, this should also be true for the other quiescent LMXBs, which do not require more energy per accreted nucleon to explain the base quiescence level and most are actually too cold to be explained using the standard model.  Therefore, we conclude that there are indeed VFXTs in globular clusters, but that they can likely not be associated with the brightest amongst the quiescent LMXB candidates ($>10^{33}$ erg s$^{-1}$) and likely also not with the slightly fainter ones ($10^{32-33}$ erg s$^{-1}$).  It is more likely that those quiescent LMXB candidates are associate with bright transients similar to the transients IGR J17480--2446 and Swift J174805.3--244637 in Terzan 5 which were as well identified previously as quiescent LMXB candidates \\citep[][]{2006ApJ...651.1098H} before they exhibited their bright outbursts \\citep[][]{2011MNRAS.412L..68D,wijnandsatel}. However, also those transients should have low duty cycles in order to be consistent with the observations of the lack of outbursts (either very-faint or bright). More and more evidence becomes available that there exist a class of transients which might have indeed very low duty cycles. This might not be unexpected because their is a strong selection effect in favor for discovering new transients with a high duty cycle.   \\subsection{Final remarks \\label{forward}}  The above discussion has demonstrated that VFXTs could be extremely faint in quiescence but some of them could reach quiescent X-ray luminosities of $10^{32}$ to even $10^{33}$ erg s$^{-1}$ depending on their accretion rate history and the amount of matter accreted. So, it would be interesting to observe quiescent VFXTs and determine their quiescent properties. Sadly, two main factors hamper the study of the thermal cooling emission of quiescent neutron-star VFXTs. First of all, many quiescent LMXBs exhibit besides the thermal component a hard, non-thermal component above 2 keV. The origin is not well understood but its presence inhibits the most accurate study of the thermal component and in some systems only this hard component can be detected. In particular, there seems a trend which indicates that the fainter a source is in quiescence, the larger the contribution is of this non-thermal component to the 0.5--10 keV X-ray luminosity \\citep[][]{2004MNRAS.354..666J}. If this trend is also valid for the quiescent state of VFXT, then this will make it difficult to maybe even impossible to put significant constraints on the thermal component in those systems.  In addition, VFXTs are difficult to discovery because their peak X-ray luminosity is below the sensitivity limits of all-sky X-ray instruments in orbit and therefore, most outbursts are missed. Often only using pointed observations of more sensitive instruments, those outbursts can be detected. But usually, those pointed observations have a very narrow field-of-view and mostly are pointed towards the Galactic center and the Galactic bulge. Consequently the interstellar absorption is usually rather high for those systems with values between $10^{22}$ cm$^{-2}$ to $10^{23}$ cm$^{-2}$. With such high column densities and the expected low surface temperature of the neutron stars in VFXTs, it is very difficult to impossible to detect the thermal component. Even more so if also the non-thermal component discussed above is present in the quiescent spectrum as well.  Despite that the situation described above looks very bleak, it might still be possible to use quiescent VFXTs as tests for the deep-crustal heating model. As described in Section~\\ref{globulars}, many X-ray transients are expected to be present in Galactic globular clusters among which many VFXTs (three sources are already currently known). Typically for many clusters the column density is not very high and the distances are typically quite accurately known. Continued monitoring of those clusters with sensitive X-ray instruments would be crucial to catch VFXTs in outburst so that later they can be studied in quiescence \\citep[][]{2009ApJ...692..584H,2010ApJ...714..894H}.  When in outburst, one specific observable property can be very useful to determine whether a VFXT is a primordial system or that it has accreted at significantly higher rate in the past: detecting millisecond X-ray pulsations. As discussed in Section~\\ref{additional} in the context of the accreting millisecond X-ray pulsar SAX J1808.4--3658, detecting such millisecond X-ray pulsations basically guarantees that the neutron star has accreted enough matter to replace the crust. Therefore, those systems cannot be primordial (hence we expect spin periods $>$10 ms for primordial VFXTs) and the neutron stars systems in those systems cannot have a hybrid crust (unless in some way the matter spins up the neutron star but is not eventually accreted on the neutron star itself). The fact that the VFXT NGC 6440 X-2 is a millisecond X-ray pulsar with a spin frequency of 4.8 ms \\citep[][]{2010ApJ...712L..58A} indicates that this source is not a primordial VFXT and that is faint quiescent emission \\citep[][]{2010ApJ...714..894H} cannot be due to non-standard heating in a hybrid neutron star crust because the crust in this system should be fully replaced and a different explanation (e.g., enhanced core cooling) is needed.  Detailed calculations have to be performed whether other observable properties during outbursts can discriminate systems with a fully accreted crust with those which have a hybrid crust (e.g., type-I X-ray bursts behavior; super-burst behavior).  Another possibility is to study those VFXTs which are active for a very long time (years to decades, instead of weeks to months), the so-called quasi-persistent sources, when their outbursts turn off. Several such quasi-persistent VFXTs have been identified \\citep[see, e.g.,][albeit that those transients are close Sgr A* and therefore have a large column density making them unsuitable for studying the soft thermal component]{2007A&A...468L..17D,2010A&A...524A..69D}. It has been found that for normal transients which accrete so long that the neutron star crust is heated considerable out of thermal equilibrium with the core and for several months to years after the end of their outburst, the observed quiescent temperature tracks the thermal evolution of the crust instead of that of the core \\citep[until thermal equilibrium is reached again; see][]{2001ApJ...560L.159W,2002ApJ...573L..45W,2003ApJ...594..952W,2004ApJ...606L..61W,2002ApJ...580..413R,2006MNRAS.372..479C,2008ApJ...687L..87C,2010ApJ...722L.137C,2010ApJ...714..270F,2011ApJ...736..162F,2011MNRAS.414L..50D,2009A&A...495..547D,2009MNRAS.396L..26D,2011MNRAS.412.1409D,2011MNRAS.418L.152D,2011A&A...528A.150D}. If a quasi-persistent VFXT can be studied right after the end of its prolonged outburst, the observed crust cooling curve could give significant insight into the heating and thermal processes which occurred in the crust during the outburst. Very likely a fully accreted crust will exhibit a different cooling profile than an hybrid crust\\footnote{We note that for the bright transient IGR J17480--2446 in Terzan 5 (which might harbor a neutron star with a hybrid crust; see Section~\\ref{additional}) the crust cooling curve has been measured \\citep[][]{2011MNRAS.414L..50D,2011MNRAS.418L.152D,Degenaar2012} and it did not meet the theoretical expectations. This could be due to a neutron star with a hybrid crust \\citep[see][for an in-depth discussion]{Degenaar2012}.}. Although this assumes that still a significant amount of heat is produced in the crust during outburst to elevate the thermal emission above the sensitivity limits of the current generation of X-ray instruments.  Obtaining crust-cooling curves for neutron stars in HMXBs (i.e., after the bright and extended type-II outbursts observed from several neutron-star Be/X-ray transients) will also be very interesting because it is expected that those curves could deviate significantly from those obtained so far for the neutron-star LMXBs. Both the effects of hybrid crusts in the neutron stars in HXMBs as well as their high magnetic field strength might produce observational effects on the crust cooling curves.  Primordial VFXTs will have different companions than systems which are only in a VFXT phase but normally they are normal transients or they used to be normal transients in the past \\citep[][]{2006MNRAS.366L..31K}. So finding the optical or IR companion star of VFXTs would be very useful to pin-pointed the most likely accretion rate history of those sources. Again globular clusters might be the best places to study this because of the low absorption for many of them compared to field targets. However, the companion star for both the primordial and the non-primordial VFXTs can be very faint \\citep[][]{2009ApJ...692..584H,2010ApJ...714..894H} and detecting them will remain a challenge, let alone obtaining spectral observations to confirm the type of companion star. E-ELT using adaptive optics might be able to detected more systems in the crowded globular clusters and it might be able to take spectra of the brightest targets.\\\\   \\noindent {\\bf Acknowledgements.}\\\\ R.W. acknowledges support from a European Research Council (ERC) starting grant. N.D. is supported by NASA through Hubble postdoctoral fellowship grant number HST-HF-51287.01-A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555.  D.P.'s work is supported by a grant from Conacyt (CB-2009/132400).  This research has made use of NASA's Astrophysics Data System. We thank Craig Heinke, Alessandro Patruno and Diego Altamirano for useful discussion. We also thank Diego Altamirano for useful comments on a previous version of this paper. We thank the organizers of the ``CompStar 2012: the physics and astrophysics of compact stars'' conference on Tahiti for organizing this nice meeting, because part of the paper was constructed there."
    },
    "1208/1208.1510_arXiv.txt": {
        "abstract": "We present the source catalog of a new \\Chandra ACIS-I observation of NGC~300 obtained as part of the \\Chandra Local Volume Survey. Our 63 ks exposure covers $\\sim88$\\% of the $D_{25}$ isophote ($R\\approx6.3$ kpc) and yields a catalog of 95 X-ray point sources detected at high significance to a limiting unabsorbed 0.35-8 keV luminosity of $\\sim10^{36}$ \\lum. Sources were cross-correlated with a previous \\XMM catalog, and we find 75 ``X-ray transient candidate'' sources that were detected by one observatory, but not the other. We derive an X-ray scale length of 1.7$\\pm$0.2 kpc and a recent star formation rate of 0.12 \\Msun yr$^{-1}$, in excellent agreement with optical observations. Deep, multi-color imaging from the {\\it Hubble Space Telescope}, covering $\\sim32$\\% of our \\Chandra field, was used to search for optical counterparts to the X-ray sources, and we have developed a new source classification scheme to determine which sources are likely X-ray binaries, supernova remnants, and background AGN candidates. Finally, we present the X-ray luminosity functions (XLFs) at different X-ray energies, and we find the total NGC~300 X-ray point source population to be consistent with other late-type galaxies hosting young stellar populations ($\\lesssim50$ Myr). We find XLF of sources associated with older stellar populations has a steeper slope than the XLF of X-ray sources coinciding with young stellar populations, consistent with theoretical predictions. ",
        "introduction": "The X-ray emission from normal, non-active galaxies originates from a mixture of stellar sources and hot, diffuse gas. In spiral galaxies, this X-ray emission is dominated by X-ray binaries (XRBs) comprising either a neutron star (NS) or a black hole (BH) accreting from a stellar companion. The fast evolutionary timescale of massive stars ($\\sim10^7$ years) makes X-ray emission from high-mass XRBs (HMXBs) nearly simultaneous with their formation \\citep{Shty+07}, while the longer-lived low-mass XRB (LMXB) systems trace older underlying stellar populations \\citep{Kong+02,Soria+02,Trudolyubov+02}; thus, the age and star formation history (SFH) of the host galaxy should correlate with the shape of the X-ray luminosity function (XLF) \\citep{Grimm+03, Belczynski+04,Eracleous+06}. Dependencies of the XRB population properties on host galaxy metallicity may provide useful constraints on binary evolution and population synthesis models.  With its excellent angular resolution ($\\sim$0\\as5) and positional accuracy, the {\\it Chandra X-ray Observatory} is the only X-ray telescope capable of separating the X-ray point source populations of nearby galaxies from diffuse emission. When combined with deep, optical {\\it Hubble Space Telescope} (\\HST) imaging, reliable source identification and optical counterpart identification may be carried out even at distances of a few Mpc. The \\Chandra Local Volume Survey (CLVS, P.I. Benjamin Williams) is a deep, volume-limited X-ray survey of five nearby galaxies (NGC~55, NGC~300, NGC~404, NGC~2403, and NGC~4214) with matched \\HST observations down to \\Mv$\\sim$0 \\citep{Dalcanton+09}. When combined with the already well-studied disks of M~31 \\citep{Kong+03} and M~33 \\citep{Williams+08,Plucinsky+08,Tullmann+11}, these galaxies contain $\\sim$99\\% of the stellar mass and $\\sim$90\\% of the recent star formation out to a distance of $\\sim$3.3 Mpc \\citep{Tikhonov+05,Maiz+02,Freedman+Madore88}. Additionally, these galaxies span a representative sample of disk galaxies with a range of masses, metallicities, and morphologies.  In this paper, we present the results of a new \\Chandra observation of the nearby SA(s)d spiral galaxy NGC~300, located at a distance of 2.0 Mpc \\citep{Dalcanton+09} in the Sculptor Group. NGC~300 is a near-optical twin of the Local Group galaxy M~33, and has a low foreground Galactic absorption of 4.09$\\times10^{20}$ cm$^{-2}$ \\citep{Kalberla+05} and is viewed at an inclination angle $i=42^{\\circ}$; Table~\\ref{galbasic} lists some basic properties of NGC~300, as well as M~33 and the Milky Way for direct comparison.  \\begin{table*}[ht] \\centering \\caption{Summary of NGC~300 Properties} \\begin{tabular}{cccc} \\hline \\hline Property\t\t& NGC~300\t& M~33\t&\tMilky Way \\\\ (1) & (2) & (3) & (4) \\\\ \\hline Distance\t\t\t& 2.0 Mpc$^a$\t\t\t\t\t& 800 kpc$^b$\t\t\t\t\t& 8.5 kpc \\\\ Type\t\t\t\t& SA(s)d$^c$\t\t\t\t\t& SA(s)cd$^c$\t\t\t\t\t& SBc \\\\ $M_B$\t\t\t& -17.66$^a$\t\t\t\t\t& -18.4$^d$\t\t\t\t\t& -20.3$^h$ \\\\ Scale length (kpc)\t& 1.3$^e$\t\t\t\t\t\t& 1.4$^e$\t\t\t\t\t\t& 3.00$\\pm$0.22$^h$ \\\\ Circular velocity \t& 97 km s$^{-1}$$^f$\t\t\t& 130 km s$^{-1}$$g$\t\t\t& 239$\\pm$5 km s$^{-1}$$^h$ \\\\ Estimated stellar mass\t& $4.3\\times10^{9}$ \\Msun$^f$\t& $4.5\\times10^{9}$ \\Msun$^g$\t& $6.43\\pm0.63\\times10^{10}$ \\Msun$^h$\\\\ \\hline \\hline \\label{galbasic} \\end{tabular} \\tablecomments{$^a$\\cite{Dalcanton+09}; $^b$\\cite{Williams+09}; $^c$NED; $^d$\\cite{VilaCostas+92}; $^e$\\cite{Munoz+07}; $^f$\\cite{Puche+90}; $^g$\\cite{Corbelli+00}; $^h$\\cite{McMillan11}} \\end{table*}  The first X-ray population studies of NGC~300 were performed with {\\it ROSAT} between 1991 and 1997 \\citep{Read+01}, identifying 29 bright X-ray sources including NGC~300 X-1, a highly variable supersoft source, and other bright sources coincident with known supernova remnants (SNRs) and \\ion{H}{2} regions. More recently, an \\XMM survey of the the X-ray source population of NGC~300 was presented by \\cite{Carpano+05} down to a limiting luminosity of $\\sim3\\times10^{35}$ \\lum in the 0.3-6 keV band. A total of 163 sources were detected in the energy range of 0.3-6 keV, and the 86 sources falling within the $D_{25}$ optical disk were further characterized using hardness ratios, X-ray fluxes, and the available ground-based optical imaging. A brief summary of previous X-ray NGC~300 studies is provided in Table~\\ref{previous_studies}.  \\begin{table*}[ht] \\centering \\caption{Summary of Prior X-ray Surveys of NGC~300} \\begin{tabular}{cccc} \\hline \\hline Observatory\t& Number of Detected \t\t& Limiting Unabsorbed Luminosity\t& References \t\\\\ & Sources (within $D_{25}$)\t& ($10^{36}$ \\lum)$^a$\t\t\t& \t\t\t\\\\ (1) \t\t\t& (2) \t\t\t\t\t& (3) \t\t\t\t\t\t& (4) \t\t\\\\ \\hline {\\it ROSAT} \t& 47 (29) \t\t\t\t\t& $\\sim4.0$ ($\\sim$7.7)\t\t\t& \\cite{Read+01} \\\\ \\XMM \t\t& 163 (86)\t \t\t\t& $\\sim0.3$ ($\\sim$0.4) \t\t\t& \\cite{Carpano+05} \\\\ \\Chandra\t\t& 95 (77)\t\t\t\t\t& 1.5\t\t\t\t\t\t\t& this work\t\t\\\\ \\hline \\hline \\end{tabular} \\tablecomments{$^a$The limiting unabsorbed luminosities correspond to the 0.1-2.4 keV and 0.3-6.0 keV energy ranges for {\\it ROSAT} and \\XMM, respectively. In parenthesis we provide expected 0.35-8 keV limiting unabsorbed luminosity assuming a power law with \\PL=1.9 and \\nH=$4.09\\times10^{20}$ cm$^{-2}$.} \\label{previous_studies} \\end{table*}  Before the publication of the full \\Chandra X-ray point source catalog, our collaboration employed this new observation of NGC~300 to study two individual objects of interest: the ``supernova impostor\" SN 2010da, which we argued is consistent with a Be-X-ray binary (BeXB, comprised of a young, rapidly rotating B-type star with strong emission lines and a NS) origin \\citep{Binder+11a}, and the BH + Wolf-Rayet binary NGC~300 X-1 \\citep{Binder+11b}. These sources are included as part of the X-ray source catalog, but are not discussed in detail in this work.  The organization of this paper is as follows: in Section~\\ref{obs} we provide the details of our new observation, data reduction procedure, and the technique used to generate the X-ray source catalog. An analysis of the X-ray source catalog, including comparisons to previous studies, hardness ratios, and the radial source distribution is given in Section~\\ref{xray_results}. In Section~\\ref{classifications}, we present a new X-ray/optical source classification scheme, identify candidate optical counterparts to our X-ray sources, and assign likely source classifications (HMXB, LMXB, SNR, or background AGN) to each source. The X-ray luminosity function, and the possible role of transient sources in determining the luminosity function shape, is discussed in Section~\\ref{XLF}. The paper concludes with a summary and the main conclusions in Section~\\ref{end}. ",
        "conclusions": "\\label{end} We have presented the X-ray point source catalog for a new, deep \\Chandra observation of NGC~300 as part of the \\Chandra Local Volume Survey. A total of 95 sources were detected in NGC~300 down to a 90\\% limiting 0.35-8 keV luminosity of $\\sim1.5\\times10^{36}$ \\lum. By comparing our X-ray source catalog to earlier observations of NGC~300 performed by \\XMM, we were able to place long-term variability constraints on each source and identify 25 XRT candidates. We fit the radial X-ray source distribution to an exponential profile and were able to derive a scale length of 1.7$\\pm$0.2 kpc, in good agreement with measurements taken at other wavelengths. We evaluated hardness ratios for each of our X-ray sources, and performed spectral fitting for those sources with more than 50 counts.  We utilized nine deep {\\it Hubble Space Telescope} fields to search for optical counterparts to each of our X-ray sources; 32 of our X-ray sources were observed in at least one of the overlapping \\HST fields. We developed a new source classification method which combined both X-ray and optical data to classify each source as a HMXB, LMXB, or background AGN, and have made our IDL routine which automates our source classification technique publicly available. We classified 47 sources as background AGN, 26 as HMXBs, 11 as LMXBs, 10 as SNRs, and one as a foreground star.  The cumulative \\lognlogs distributions were presented in both 0.5-2 and 2-8 keV bands and found to be consistent with those of other spiral galaxies with a dominant HMXB population. Additionally, we divided our X-ray sources into two radial bins and modeled the radially-resolved \\lognlogs distributions with power laws. Using the SFHs of \\cite{Gogarten+10}, we were able to assign an average to the underlying stellar populations, and we find the measured XLF slopes and stellar population ages to be in good agreement with model predictions. The remaining four CLVS galaxies (NGC~55, NGC~404, NGC~2403, and NGC~4214) will provide additional constraints on XLF-age models with upcoming works."
    },
    "1208/1208.4103_arXiv.txt": {
        "abstract": "We present a catalog of 1964 isolated, compact neutral hydrogen clouds from the Galactic Arecibo L-Band Feed Array Survey Data Release One (GALFA-HI DR1). The clouds were identified by a custom machine-vision algorithm utilizing Difference of Gaussian kernels to search for clouds smaller than 20\\arcmin. The clouds have velocities typically between $|$\\vlsr$|$ $= 20 - 400$ \\kms, linewidths of $2.5-35$ \\kms, and column densities ranging from $1-35 \\times 10^{18}$ cm$^{-2}$. The distances to the clouds in this catalog may cover several orders of magnitude, so the masses may range from less than a Solar mass for clouds within the Galactic disc, to greater than $10^4$ M$_\\odot$ for HVCs at the tip of the Magellanic Stream. To search for trends, we separate the catalog into five populations based on position, velocity, and linewidth: high velocity clouds (HVCs); galaxy candidates; cold low velocity clouds (LVCs); warm, low positive-velocity clouds in the third Galactic Quadrant; and the remaining warm LVCs. The observed HVCs are found to be associated with previously-identified HVC complexes. We do not observe a large population of isolated clouds at high velocities as some models predict. We see evidence for distinct histories at low velocities in detecting populations of clouds corotating with the Galactic disc and a set of clouds that is not corotating. ",
        "introduction": "The recent discovery of compact clouds over the entire velocity range of our Galaxy and the Local Group has posed interesting questions as to their role in galaxy evolution and the Galactic ISM \\citep[e.g.][]{lockman02, ford08, ryanweber08, benbekhti09, stanimirovic08, heitsch09, begum10, giovanelli10, hsu11}. Despite the similarity of the individual clouds when observed with the hyperfine neutral hydrogen 21 cm line, different classes of clouds have been identified and many possible origins have been proposed. The detection of stellar outflows in HI points to the possibility that some small clouds could be produced by evolved stars \\citep[e.g.][]{matthews11}. At the disc-halo interface, a population of clouds has been observed rotating with the Galaxy that may have been pushed off the disc by stellar feedback \\citep{ford08, ford10}. In the Galactic halo, compact clouds may represent the initial cooling seeds from the multiphase halo medium that will eventually fuel the disk \\citep{maller04, joung11}, or the remnants of larger complexes \\citep{heitsch09}. The continual production of these clouds could provide a significant source of the fuel for the Milky Way's ongoing star formation. The abundance of newly discovered small Local Group dwarf galaxies \\citep[and others]{belokurov2006,zucker2006,irwin07}, one of which is particularly gas-rich \\citep{ryanweber08}, raises the possibility that some of these clouds represent previously undiscovered dwarf galaxies. It is even possible that some compact clouds are outside the halo and physically large \\citep{giovanelli10}.  Unraveling the nature of the compact clouds requires a large, well-defined sample. Historically, finding such a sample has been difficult due to the limited area, sensitivity, and resolution covered by HI surveys, as well as the difficulty of recovering compact objects confused by diffuse Galactic emission \\citep{stil06, ford08}.    The latter problem is evident in the lack of detections of halo clouds (commonly referred to as high-velocity clouds; HVCs) between \\vlsr $= -90$ to 90 \\kms \\citep{wakker97,peek09}, and the recent, optical discovery of gas-rich dwarf satellites at velocities where Galactic emission is prevalent \\citep{irwin07}. A large, sensitive survey with sufficient kinematic and spatial resolution to recognize clouds that are near Galactic emission is needed.  The GALFA-HI (Galactic Arecibo L-Band Feed Array HI) Survey \\citep{peek11} provides the data needed to detect a large sample of compact clouds across the range of velocities of the Galaxy and Local Group. At completion it will cover 13,000 deg$^2$ (58\\% of which is available in Data Release 1\\footnotemark and cataloged here; DR1) \\footnotetext{The GALFA-HI DR1 is publicly available at http://purcell.ssl.berekley.edu} over -1\\deg $< \\delta <$ 38\\deg~and the entire RA range between \\vlsr $=\\pm650$ \\kms. The GALFA-HI survey has a spatial resolution of 4\\arcmin~and a channel spacing of 0.184 \\kms. The sensitivity at 6$\\sigma$ to compact, warm (FWHM = 15 \\kms) HI clouds corresponds to masses as low as 7 $\\times$ 10$^{-2}$ M$_\\odot$ at 1 kpc and 7 $\\times$ 10$^{4}$ M$_\\odot$ at 1 Mpc (see \\S \\ref{sec:obs}). As the sensitivity of a survey to compact, spectrally resolved sources scales inversely with both the antenna beam solid angle and the noise for a given channel width, GALFA-HI is uniquely positioned among other large-area Galactic HI surveys. With the sensitivity of a single-dish radio telescope and resolution approaching that of a compact array, GALFA-HI is 2.5 to 80 times more sensitive to compact clouds than other large-area surveys (2.5: IGPS; \\citealt{taylor03,stil06b,mccluregriffiths05}, 15:GASS/ EBHIS; \\citealt{mccluregriffiths09,winkel10}, 80: LAB; \\citealt{kalberla05}). Initial investigations of GALFA-HI data identified many compact clouds, both visually \\citep{begum10, stanimirovic06} and through an automated algorithm \\citep{hsu11}.  In this paper we present a catalog of compact clouds (4\\arcmin--20\\arcmin) from the GALFA-HI Survey DR1 that has been created using a new technique designed to recover faint clouds and identify structures near diffuse Galactic emission. The methods we use are detailed in \\S\\ref{sec:method}, with the completeness of the catalog assessed in \\S \\ref{sec:sense} and the catalog presented in \\S\\ref{sec:properties}.   The properties of the clouds cataloged are described in \\S \\ref{sec:results} where the clouds are separated into five populations.  A discussion of the implications of our results can be found in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this section we discuss how the different populations of clouds we have defined in \\S\\ref{sec:results} fit into current theories and how they relate to previous observations.  \\subsection{UCHVCs}  UCHVCs are especially of interest because it has been suggested that small HVCs may represent physically large structures at megaparsec distances, much farther than the classical complexes\\citep{braun99, giovanelli10}. It is very hard to obtain direct distance information for such small clouds, given the unlikely chance of an overlap with a halo star;  however, if the UCHVCs reside near larger complexes in position-velocity space, we expect they are associated with these relatively nearby structures, rather than being independent clouds at much larger distances.  As discussed in \\S\\ref{sec:results_hvcs}, all but a few UCHVCs have $D<25\\degr$ to an HVC in the WvW catalog. At negative velocities, this leaves 86\\% of the available position-velocity space with $-400 <$\\vlsr$<-90$~\\kms~empty. The presence of UCHVCs found only near other known HVC complexes is consistent with the bulk of the Galactic hot halo not cooling due to linear thermal instabilities \\citep{binney09, joung11}.  If the bulk of the halo were thermally unstable to linear perturbations, we would expect small neutral condensations to appear all over phase space, rather than only concentrated toward known, larger HVCs. This result instead supports the idea that some HVCs may be seeded by larger perturbations in the halo, such as filamentary streams of gas impinging on the Galaxy from the intergalactic medium or satellite galaxies\\citep{keres09, joung11}.  If UCHVCs are indeed associated with larger clouds, they are most likely at similar distances. The median UCHVC properties indicate a mass of $\\sim200$M$_{\\odot}$ and a physical size of $\\sim10$pc at a distance of 10kpc and a mass of $\\sim2\\times10^4$M$_{\\odot}$ and a physical size of $\\sim100$pc at 100kpc.  These distances were chosen based on the distance constraints to many HVC complexes \\citep{thom06,wakker01, wakker08}, and the likely distance of the tail of the Magellanic Stream given simulations \\citep{besla10, connors06}.  The mass of these clouds indicate they are not likely to survive the trip to the Galactic disk and will instead become part of the multi-phase Galactic halo \\citep{heitsch09}.  The exception would be if some of the UCHVCs are embedded within extended, highly ionized shells.  \\subsubsection{Warm, Positive-Velocity Q3 Clouds}\\label{sec:warmq3discussion}  In \\S\\ref{sec:warmq3results} we state that the vast majority of the positive velocity warm low velocity clouds ($0<$\\vlsr$<90$~\\kms) are located in the third Galactic quadrant and are likely to be associated with the W HVC complexes. In Figure \\ref{fig:ra_dec} we have plotted the \\vlsr~ velocities corresponding to \\vgsr$=0$~\\kms~as a curved line. Clouds on this line could be at rest relative to Galactic rotation. The majority of the positive-velocity warm clouds in Q3 are bounded by the \\vgsr$=0$~\\kms~line indicating these clouds may be infalling clouds with positive LSR velocities due to Galactic rotation.  The HVC complexes WA, WB, and WC (e.g. WvW) are in the same region of sky as the warm Q3 clouds, and if we extend the associations with HVCs from \\S\\ref{sec:results_hvcs} below 90 \\kms, many of the Q3 clouds would be considered part of these complexes. \\cite{thom06} measured a direct distance constraint to clouds in the W complexes and found a distance of $\\sim9$ kpc.  This results in a z-height of 7 kpc and galactocentric distance of 12 kpc. We consider three scenarios for a cloud at this distance at ($\\alpha$,$\\delta$)=(150\\deg,15\\deg):  a non-rotating cloud, a co-rotating cloud, and a co-rotating cloud with a vertical lag. For a non-rotating cloud (\\vgsr=0), we would observe \\vlsr $\\simeq95$~\\kms due to the rotation of the Galaxy. If the cloud were co-rotating at 220~\\kms, we would observe \\vlsr$\\simeq10$~\\kms. If the cloud were co-rotating, but with a vertical lag of 20~\\kms~kpc$^{-1}$, we would observe \\vlsr $\\simeq65$~\\kms. All of these calculations are without any vertical infall which would lower the observed \\vlsr. The velocities we observe for the warm Q3 clouds are consistent with non-rotating clouds with infalling velocities of $\\sim$10-70~\\kms~ or co-rotating clouds in a lagging halo. The warm Q3 clouds may fall under the category of the `low-velocity halo clouds' \\citep{peek09}, and may also represent a bridge connecting the HVC W complexes to the Galactic disc.    \\subsection{Cold Low-Velocity Clouds} The cold, low-velocity clouds are certainly associated with the disc of the Galaxy. There are greater numbers of these clouds at lower velocities and there is no apparent relationship to the GSR frame like for the warm Q3 clouds (\\S\\ref{sec:warmq3discussion}). The cold LVCs may be related to the clouds studied by \\cite{ford08, ford10} who identified a population of discrete, cold clouds 600-1000pc out of the plane of the Galaxy that are co-rotating with the disc. Their measurements were made on a set of clouds with Galactic latitude less than $20\\degr$ and at a distance of $\\sim$8kpc in the first and third Galactic quadrants. Ford et al. determined distances by selecting clouds above the tangent point velocity, which is the maximum velocity allowed by Galactic rotation toward sightlines of the inner Galactic plane. The GALFA-HI DR1 does not include the inner Galactic plane so we cannot use the tangent point method to determine distances. Much of this catalog is at high Galactic latitudes where the Ford clouds would be closer ($<1$kpc) and therefore appear substantially larger, up to two degrees, though many of their clouds were unresolved. The masses of the Ford sample range from $\\sim100$ to $\\sim5000$ M$_\\odot$, while the cold LVCs from our catalog would have masses of $\\sim1$ M$_\\odot$ at a similar vertical position. High latitude analogs of the Ford clouds would not be included in this catalog due our size limit of 20\\arcmin, while analogs to the clouds in this catalog would not have been detected by Ford et al. due to the small sizes and masses. However, the clouds in the Ford sample have a median linewidth of 10.6 \\kms, while our cold clouds have a median linewidth of  8.4 \\kms, indicating both populations have typical temperatures less than 5000K. The cold LVCs in this catalog are not physically the same as the Ford clouds, but may have similar temperature, vertical structure, and kinematics.  Two possible origins of the cold LVCs are a Galactic fountain process or recooling disc-halo interface gas. In a Galactic fountain model, gas is expelled out of the plane of the Galaxy by some energetic process and then rains back down out of the halo \\citep{shapiro76}. Whether the ejected gas is localized to regions of star formation, or rapidly blends with existing warm/hot lower halo gas and condenses later is still unknown.   This will be investigated further with models and a larger population of GALFA-HI clouds in the future. Regarding recooling gas, \\cite{heitsch09} observed small clouds forming from the warm remnants of disrupted HVCs in their simulations as the clouds became buoyant close to the disk. The large number of positive-velocity cold LVCs requires some process to drive them away from the disc, be it the buoyancy of accreting gas or feedback from star formation in the disc.   \\subsection{Warm Low-Velocity Clouds}  After classifying the positive-velocity warm clouds in the third Galactic quadrant as possible low-velocity halo clouds, we see that 90\\% of the remaining warm low-velocity clouds have negative velocities. This agrees well with previous studies of intermediate velocity ($|$\\vlsr$|$ $=40-90$ \\kms) clouds that found most of this gas is at negative velocities and has a vertical distance ranging from 0.5 to 3 kpc (See \\cite{albert04} for an extensive discussion). The warm low-velocity clouds detected here may be associated with IVCs. If these clouds are within 3kpc of the disc, they would have masses of $\\sim1-100$M$_\\odot$. The origin of the clouds is uncertain but the higher metallicities for some of the large IVCs \\citep{wakker01} suggests a relation to processes within the Galactic disc. As shown in Figure \\ref{fig:histo2}, the velocity and line width distributions are continuous across the 90 \\kms~ and 15 \\kms~ boundaries, respectively. This suggests that the warm LVCs may be partially low-velocity HVCs and partially warm or turbulent cold LVCs. This is different from the cold LVCs that appear separate from the HVCs in velocity. From this reasoning, some of the warm LVCs are most likely the remnants of disrupted HVCs, while other warm LVCs may have come from, or have yet to become, cold LVCs.  In any case, the negative velocities of these clouds indicate they will be future star formation fuel."
    },
    "1208/1208.6330_arXiv.txt": {
        "abstract": "The Gas Pixel Detector was designed and built as a focal plane instrument for X-ray polarimetry of celestial sources, the last unexplored subtopics of X-ray astronomy. It promises to perform detailed and sensitive measurements resolving extended sources and detecting polarization in faint sources in crowded fields at the focus of telescopes of good angular resolution. Its polarimetric and spectral capability were already studied in earlier works. Here we investigate for the first time, with both laboratory measurements and Monte Carlo simulations, its imaging properties to confirm its unique capability to carry out imaging spectral-polarimetry in future X-ray missions. ",
        "introduction": "\\label{Intro}  The Gas Pixel Detector (GPD) \\cite{2001Natur.411..662C, 2006NIMPA.566..552B, 2007NIMPA.579..853B} has been designed and built to perform time resolved imaging spectral-polarimetry of X-ray celestial sources by means of the photoelectric effect. It is the long sought quantum leap in sensitivity with respect to the classical techniques \\cite{2003NIMPA.510..170S}. Successively the same method was implemented exploiting an alternative technique without any 2-D imaging capability \\cite{2007NIMPA.581..755B}.  While the polarimetric and the spectroscopic performances of the GPD have been already studied in detail \\cite{2008NIMPA.584..149M, 2010NIMPA.620..285M}, its imaging properties have been estimated, so far, only by Monte Carlo simulations evaluating the point of conversion of the impinging photon with a suitable algorithm \\cite{2003SPIE.4843..383B}. Still the GPD was devised primarily to be exploited at the focus of conventional and multi-layer hard X-ray optics \\cite{2006SPIE.6266E..85M}. At this purpose the GPD has been considered for different proposed space missions \\cite{2008SPIE.7011E..62S,2010ExA....28..137C,2010xpnw.book..269B,2010SPIE.7732E..38S}. While X-ray optics with the exquisite quality of those of Chandra will not be again available in the near future, for POLARIX (or XIPE, proposed in 2012 to ESA as a small mission for a launch on 2017) the JET-X optics (three complete mirrors, two Flight Models, FM, the heritage of the former Spectrum X-Gamma mission, and one Qualification Model), can still be used. They have a Point Spread Function (PSF) with a measured Half Energy Width (HEW) at 1.5-keV of 15.l'' (FM-1) and 14.6'' (FM-2) \\cite{1997SPIE.3114..392W}.  We report on the first measurement, with a narrow beam at 4.5 keV, of the position resolution of a GPD filled with a mixture of He-DME 80-20 (1 bar) with a large Gas Electron Multiplier (GEM), This proves the capability of the GPD to perform space-resolved measurements, such as those of Pulsar-Wind Nebulae, Supernova Remnants or X-ray jets with a suitable accuracy, and to observe the faintest source accessible to X-ray polarimetry especially in crowded fields \\cite{2012SPIE.8443S }. We also compared such measurement with the Monte Carlo simulation relating the characteristics of the image to the modulation factor. Finally we showed the effect on imaging of the non-uniformity of the drift electric field in case the GEM plane, and in particularly the guard-ring, is not powered. ",
        "conclusions": "\\label{Conclusions.} We studied the position resolution of the GPD by means of laboratory measurements and Monte Carlo simulations, at 4.5 keV, using a narrow X-ray beam. After subtracting the contribution of the beam size which was measured by means of a scanning technique, we obtained a position resolution more than twice better than that already found with a Micro Pattern Gas Detector with 200 $\\mu$m pitch and a GEM with a larger pitch \\cite{2003SPIE.4843..383B}. The measured position resolution is very close to the Monte Carlo results. By using the latter and equation \\ref{bidimgauss_symmetric_int}, the Half Energy Width (HEW, spatial) found is 36.3 $\\mu$m. The expected HEW  (spatial) by the JET-X optics (with a focal length of 3500 mm and a HEW (angular) of 19.3'' \\cite{2010xpnw.book...79L}), is 327 $\\mu$m, a factor of nine larger, including also the effects of inclined penetration of the photons collected in the absorption depth of the GPD. The angular resolution of an experiment with the GPD at the focus of such X-ray optics is, therefore, driven primarily by the quality of the latter.  We detected the presence of wings in the image of the beam in both the real data and the Monte Carlo data. Such wings were studied primarily by simulations and are due to an inversion in the location of the impact point due to peculiar event trajectories and to the blurring. The core events have a larger modulation factor with respect to the events in the wings suggesting that a spatial selection of the data, a different weight, or a better pattern recognition, especially for polarimetry of point-like celestial sources, would lead to an improved statistical significance.  Finally we showed the effect of the presence of the large guard-ring on the image profile of the narrow beam."
    },
    "1208/1208.2039_arXiv.txt": {
        "abstract": "{Moving magnetic features (MMFs) are small-size magnetic elements that are seen to stream out from sunspots, generally during their decay phase. Several observational results presented in the literature suggest them to be closely related to magnetic filaments that extend from the penumbra of the parent spot. Nevertheless,  few observations of MMFs streaming out from spots without penumbra have been reported. The literature still lacks of analyses of the physical properties of these features.} {We investigate physical properties of monopolar MMFs observed around a small pore that had developed  penumbra in the days preceding our observations and compare our results with those reported in the literature for features observed around sunspots. } {We analyzed NOAA 11005 during its decay phase with data acquired at the Dunn Solar Telescope in the \\ion{Fe}{i}~$617.3$~nm and the \\ion{Ca}{ii}~$854.2$~nm spectral lines with IBIS, and in the G-band. The field of view showed monopolar MMFs of both polarities streaming out from the leading negative polarity pore of the observed active region. Combining different analyses of the data, we investigated the temporal evolution of the relevant physical quantities associated with the MMFs as well as the photospheric and chromospheric signatures of these features.} {We show that the characteristics of the investigated MMFs agree with those reported in the literature for MMFs that stream out from spots with penumbrae. Moreover, observations of at least two of the observed features suggest them to be manifestations of emerging magnetic arches. } ",
        "introduction": "\\label{intro} Moving magnetic features (MMFs) are small-size magnetic elements seen to stream out from sunspots, especially during their decay phase. These features can appear as isolated features of polarity the same or opposite with respect to the spot they stream from (monopolar), or as couples of  features of opposite polarities (dipolar). Their horizontal velocities are usually below 1~km~s$^{-1}$ for  dipolar and monopolar features with same polarity as the spot, while higher velocities have been reported for monopolar features with opposite polarity with respect to the spot.  Recent observations have shown that MMFs form in the  middle and outer penumbrae of sunspots and that they stream out following paths traced by horizontal penumbral filaments (e.g., Bonet et al. \\cite{bonet2004}, Sainz Dalda \\&  Mart\\'{i}nez Pillet \\cite{sainz-pillet}, Ravindra \\cite{ravindra2006}).  % This has led to the hypothesis that they might be generated by Evershed flows (e.g., Cabrera Solana et al. \\cite{cabsol2006}). More recently, by analyzing spectropolarimetric data, Kubo et al. (\\cite{kubo2007}, \\cite{kubo2007_2}) found that MMFs located along lines extrapolated from the horizontal component of the sunspot magnetic field display different properties from MMFs located along lines extrapolated from the vertical component.  Some MMFs, especially monopolar ones, are associated with bright features observed at various layers of the solar atmosphere (Harvey \\& Harvey  \\cite{harvey1973}, Hagenaar \\& Frank \\cite{hagenaar2006}, Lin et al. \\cite{lin2006}, Choudhary \\& Balasubramaniam \\cite{choudhary2007}). Eruptive events, such as Ellerman Bombs (Nindos \\& Zirin \\cite{nindos1998}, Socas Navarro et al. \\cite{socasnavarro2006}), microflares (Kano et al. \\cite{kano2010}) or even CMEs (Zhang \\& Wang \\cite{zhang2002}), have also been reported in correspondance with MMFs. In addition, these features have recently been associated with shock waves traveling outward in the inner atmosphere (Lin et al. \\cite{lin2006}, Socas Navarro et al. \\cite{socasnavarro2006}, Ruytova \\& Hagenaar \\cite{ryutova2007}).  Dipolar features have been interpreted as manifestations of $\\Omega$ ( e.g., Schlichenmaier \\cite{schlichenmaier2002}, Weiss et al. \\cite{weiss04}) or U loops (e.g., Zhang et al. \\cite{zhang2003}, Sainz Dalda $\\&$ Bellot Rubio \\cite{sainzdalda2008}, Kitiashvili et al. \\cite{kitiashvili2010}) that travel along penumbral filaments.  Weiss et al. (\\cite{weiss04}) have proposed that monopolar features of the same polarity of the parent spot are flux tubes severed from the magnetic field of the spot. Those authors also interpreted monopolar features of opposite polarity with respect to the parent spot  as flux tubes that emerge in the photosphere at shallow angles. In addition, MMFs have been proposed to be manifestations of shocks (monopolar features) and solitons (dipolar features) traveling along the penumbral filaments (Ryutova \\& Hagenaar \\cite{ryutova2007}) .  All these interpretations assume the presence of a penumbra or inclined magnetic filaments around the spot generating the MMFs. Besides, some authors have explicitly pointed out the need for a visible penumbra to generate these features (Mart\\'{i}nez Pillet \\cite{martinezpillet2002}, Zhang et a. \\cite{zhang2003}). Nevertheless, MMFs have also been observed to stream out from spots without a penumbra. In particular, Harvey \\& Harvey (\\cite{harvey1973}) evinced from a statistical analysis of magnetograms that MMFs are associated with the presence of  moat regions rather than with penumbrae. Zuccarello et al. (\\cite{zuccarello2009}) observed MMFs streaming out from the naked spot \\footnote{A spot that had developed and then lost a visible penumbra during its evolution.} of an active region during its decay phase. Consistent with Harvey \\& Harvey (\\cite{harvey1973}), these authors also reported evidence of moat flows, whose connection with  penumbral filaments  had been debated in the past (e.g., Vargas et al. \\cite{vargas2008}, or Deng et al. \\cite{deng2007}) and whose independence from the presence of inclined components of the magnetic field had been recently demonstrated by magnetic hydro dynamic (MHD) simulations (Rempel \\cite{rempel2011}). This independence was also recently confirmed by Verma et al. (\\cite{verma2012}), who observed MMFs of various types streaming from a naked spot and from a pore surrounded by moat flows.  In this paper, we report on the first high cadence (52~s), high spatial resolution ($\\simeq~0.4$~arcsec), full-Stokes observations of monopolar MMFs observed around a pore with a light bridge (for a description of the photospheric structure of a typical pore with a light bridge, see Giordano et al. \\cite {giordano2008}) surrounded by a moat flow. This feature had developed a visible penumbra during days preceding our observations and should be therefore classified as naked spot. Nevertheless, for consistency with previous studies that have investigated the physical properties of this same feature (Stangalini et al. \\cite{Stangalini2012}, Sobotka et al. \\cite{sobotka2012}), we will address it as a pore  in the following. This choice is also motivated by the fact that the magnetic field of the structure does not exceed 2000 G (see Sec. \\S 4), so that it can be classified as a pore according to the classification of Zwaan (\\cite{Zwaan1987}). The paper is organized as follows: in \\S~\\ref{obs} and \\S~\\ref{ana} we describe the observations and computation of the relevant physical quantities;  in Section~\\ref{evo} we describe the evolution of the observed active region and the detection of the investigated MMFs; in Section~\\ref{res} we present our results; in Section~\\ref{disconc} we discuss and compare our findings with the recent literature; conclusions are drawn in Section~\\ref{sec:Conclusions}. ",
        "conclusions": "\\label{sec:Conclusions} We analyzed IBIS spectropolarimetric data of monopolar MMFs observed around a pore that had lost its penumbra. From the evolution of six tracked MMFs, we found that they share the following characteristics:  \\begin{itemize} \\item They move in a moat flow. \\item Type P features are generated in a region where the magnetic field surrounding the pore is highly inclined, whereas type N features stream from a region where the field surrounding the pore is vertical. \\item Type P are bigger than type N features, but less spatially coherent. \\item Type P features move at average horizontal speeds of $\\approx$ 1 km~s$^{-1}$,  while type N features are slower (hundreds of m~s$^{-1}$); both types decelerate with the increase of the distance from the pore. \\item LOS plasma motions within the MMFs are mostly upflow for type P and downflow for type N features. \\item For type P features, the appearance of bright structures in photospheric images occurs approximately 15 minutes after the formation of the MMFs. For type N features, bright structures in photospheric images are observed before their detachment from the pore field. \\item Variable brightenings in the \\ion{Ca}{ii} core frames are observed in correspondence of type N features. \\item The absolute value of the magnetic flux associated with MMFs is within 100 and 250~G. \\item The magnetic field strength distribution ranges from 500 G to 1700 G. The most likely values are 1.3 and 1.5 kG for type P and type N features, respectively. \\item The inclination of the magnetic field varies with time, although the most likely configuration is almost vertical for both types of features, with smaller inclination found for type N features. \\end{itemize}  All these properties are commensurate with the properties of monopolar MMFs streaming from spots with penumbrae. We conclude that they are very likely the manifestation of the same physical processes (as  also recently suggested by Sainz Dalda et al. \\cite{sainz2012}).  Ultimately, the main peculiarity of our observations with respect to previous studies is the high temporal cadence ($52$~s) of the data employed, which allowed us to investigate the temporal evolution of the physical properties of the MMFs observed. The evolution of the two type N features closest to the pore is driven by the surrounding environment, as these MMFs are always embedded in the outskirts of the parent pore, where new flux emerges or is expelled. For these features,  the appearance of the photospheric brightening precedes the formation of the MMF. The MMF farthest from the pore travels away in a more regular trajectory. These results suggest that type N features are magnetic flux bundles expelled from the boundary of the pore. During the evolution of type P MMFs, we found instead a slow decrease of both the magnetic flux and the upward LOS plasma motions, while the field becomes more vertical the larger the distance travelled from the pore. The photospheric brightening is delayed with respect to the appearance of the MMFs. Similar observational evidence has been recently interpreted as magnetic arches that rise from the photosphere to the higher layers of the atmosphere (Mart\\'{i}nez Gonz\\'{a}lez \\& Bellot Rubio \\cite{martinez2009}), thus suggesting that type P MMFs might be the manifestation of a similar physical process.  Similar interpretations on the nature of monopolar MMFs have been already suggested (e.g., Weiss et al. \\cite{weiss04}, Kubo et al. \\cite{kubo2007}). Results from numerical simulations that can confirm these interpretations, as well as explain in general the physical nature of the various types of MMFs  and reproduce their observed physical properties have not yet been presented in the literature."
    },
    "1208/1208.2513_arXiv.txt": {
        "abstract": "The Campanelli-Lousto solutions of Brans-Dicke theory, usually reported as black holes  are reconsidered and shown to describe, according to the values of a parameter, wormholes or naked singularities. The veiled Schwarzschild metric recently used as an example to discuss conformal frames and their equivalence corresponds to a special case of the CL metric. The conformal cousins of these solutions, and of the Riegert black hole solution of conformally invariant Weyl theory, are analysed, leading to a word of caution when interpreting physically spacetimes generated via conformal transformations from known seed solutions. ",
        "introduction": "Introduction}   There are relatively few exact solutions of alternative theories of gravity, although many such theories are currently studied with various motivations ranging from the possibility of using them to explain the current acceleration of the universe without dark energy, or for their properties in the early universe, as low-energy effective theories for quantum gravity, as emergent gravity theories, or just as toy models to understand which properties are, or are not, desirable in a theory of gravity (for recent reviews see \\cite{CliftonPadillaetc, SotiriouFaraoni, DeFeliceTsujikawa}).  When getting to know a theory of gravity, it is important to understand its spherically symmetric solutions and especially its black holes. The protoype alternative to Einstein's theory of General Relativity (GR) was Brans-Dicke theory \\cite{BD}, later generalized to scalar-tensor gravity \\cite{ST}. An early classification of spherical solutions of Brans-Dicke theory was given by Brans \\cite{Bransspherical}. A common tool used in scalar-tensor and other theories of gravity is that of conformal transformations, which relate the physics in one conformal representation of the theory (``Jordan frame'') to another (``Einstein frame''). Conformal transformations are useful to generate solutions of a theory from known seed solutions of another, but also to relate different solutions within the same theory. There has been a debate about the physical equivalence of conformal frames (see, {\\em e.g.}, \\cite{Flanagan, FaraoniNadeau} and the references therein) and recently the ``veiled Schwarzschild spacetime'' has been used as an example for discussions of conformal frames \\cite{DeruelleSasaki, ValerioAlex, AlexFirouzjaee}.    Among the known spherically symmetric and static solutions of Brans-Dicke theory are the Campanelli-Lousto spacetimes \\cite{CampanelliLousto}, which are usually reported as black holes or ``cold black holes'' (because they have zero surface gravity and temperature \\cite{coldBHs}). It turns out that the veiled Schwarzschild spacetime is a special case of the Campanelli-Lousto metrics corresponding to certain fixed values of the parameters. Given the widespread use of conformal transformations in gravity and in cosmology, one would like to understand better the veiled Schwarzschild metric and the more general Campanelli-Lousto class to which it belongs, as well as veiled black holes in other theories of gravity.  Extra motivation is provided by the finding that generalized Brans-Dicke solutions describe asymptotically Lifschitz black holes in the Jordan (but not in the Einstein) frame \\cite{asympLifschitz}.  In this paper we will first show (in Sec.~\\ref{section2}) that the three-parameter Campanelli-Lousto class of solutions of Brans-Dicke theory does not describe black holes. It corresponds, instead, to wormholes or naked singularities, respectively, according to the value of one of the parameters. We then identify, in Sec.~\\ref{section3}, the veiled Schwarzschild spacetime with a special case of the Campanelli-Lousto class and discuss its properties. A similar analysis is performed in Sec.~\\ref{section4} for the veiled Riegert black hole, which is a solution of Weyl's theory of gravity. It is found that caution is needed not to confuse Einstein frame metrics with scaling units of mass, length, and time with their versions with fixed units, and that special care must be taken when interpreting physically even straightforward mathematical results (Sec.~\\ref{section5}). Sec.~\\ref{section6} contains the conclusions. ",
        "conclusions": ""
    },
    "1208/1208.3045_arXiv.txt": {
        "abstract": "An accurate picture of how free electrons are distributed throughout the Milky Way leads to more reliable distances for pulsars, and more accurate maps of the magnetic field distribution in the Milky Way. In this paper we test 8 models of the free electron distribution in the Milky Way that have been published previously, and we introduce 4 additional models that explore the parameter space of possible models further. These new models consist of a simple exponential thick disk model, and updated versions of the models by Taylor \\& Cordes and Cordes \\& Lazio with more extended thick disks. The final model we introduce uses the observed H$\\alpha$ intensity as a proxy for the total electron column density, also known as the dispersion measure (DM). Since accurate maps of H$\\alpha$ intensity are now available, this final model can in theory outperform the other models. We use the latest available data sets of pulsars with accurate distances (through parallax measurements or association with globular clusters) to optimise the parameters in these models. In the process of fitting a new scale height for the thick disk in the model by Cordes \\& Lazio we discuss why this thick disk cannot be replaced by the thick disk that Gaensler et al. advocated in a recent paper. In the second part of our paper we test how well the different models can predict the DMs of these pulsars at known distances. We base our test on the ratios between the modelled and observed DMs, rather than on absolute deviations, and we identify systematic deviations between the modelled and observed DMs for the different models. For almost all models the ratio between the predicted and the observed DM cannot be described very well by a Gaussian distribution. We therefore calculate the deviations N between the modelled and observed DMs instead, and compare the cumulative distributions of N for the different models. Almost all models perform well, in that they predict DMs within a factor of 1.5--2 of the observed DMs for about 75\\% of the lines of sight. This is somewhat surprising since the models we tested range from very simple models that only contain a single exponential thick disk to very complex models like the model by Cordes \\& Lazio. We show that the model by Taylor \\& Cordes that we updated with a more extended thick disk consistently performs better than the other models we tested. Finally, we analyse which sightlines have DMs that prove difficult to predict by most models, which indicates the presence of local features in the ISM between us and the pulsar. ",
        "introduction": "Knowledge of the distribution of the free electrons in the Galaxy, $n_{\\mathrm{e}}$, is important for several reasons. First, the detailed structure of $n_{\\mathrm{e}}$ is an essential ingredient for a full understanding of the multi-phase character of the Galactic interstellar medium (ISM). Furthermore, a model for the distribution of the thermal electrons is a crucial tool for estimating the distances of pulsars from their observed dispersion measure (DM)\\footnote{The dispersion measure is the integral of the free electron density along the line of sight, and is expressed in units of cm$^{-3}$pc.}. Finally, through its integral DM, an $n_{\\mathrm{e}}$-model is required for an analysis of the strength and structure of the Galactic magnetic field from rotation measures (RM), because RM is the integral over the line of sight of the product of the l.o.s.-component of the magnetic field and $n_{\\mathrm{e}}$ \\citep[see, e.g.,][]{han2006, brown2007, men2008, noutsos2008, roy2008, vallee2008, nota2010}.   In our paper we analyse 8 models of the Galactic free electron density $n_{\\mathrm{e}}$ that have appeared in the literature. These models range from single exponential thick disks, see, e.g., \\citet{berkhuijsen2006}, \\citet{berkhuijsen2008} and \\citet{gaensler2008}, models that consist of a thin and a thick disk with different radial and vertical scale heights \\citep{gomez2001}, and the complex, multi-component models by \\citet{taylor1993}, or `TC93', and \\citet{cordes2002}, also known as NE2001, that consist of an axisymmetric thin and thick disk, spiral arms, and a contribution by the Gum nebula (and, in the case of NE2001 additional components). We add four models to this list. The first model uses an exponential thick disk that is fitted to the latest available data on pulsars with accurate distances. We used a similar approach to fit new scale heights of the thick disk component in TC93 and NE2001, whilst maintaining the contributions by the other components in these models.  With our final model we consider descriptions in which the prediction of DM is based on another observable quantity, viz. emission measure (EM)\\footnote{The emission measure is defined as the line-of-sight integral of the square of the free electron density, and is expressed in units of cm$^{-6}$pc.}. We then test the predictive quality of the available models using a new formalism that we developed.  We arranged this paper as follows. We present and discuss the data in Sect. \\ref{s-data}, and in Sect. \\ref{s-model} we give an overview of the different models that we test. In Sect. \\ref{s-details} we describe the practicalities of modelling DM, in particular of the four models we propose, and our criteria for selecting usable sightlines. We then test the predictive quality of the available models in Sect. \\ref{s-comp}. Finally, in Sect.~\\ref{s-anomalousn} we discuss sightlines for which the DMs prove very difficult to predict. Throughout this paper we have calculated statistics in a way that is robust against outliers. ",
        "conclusions": "In this paper we have analysed how well twelve different models of the Galactic free electron density $n_{\\mathrm{e}}$ predict the measured DMs of pulsars at well-known distances. These distances were derived from either parallax measurements, or from association of the pulsar with a globular cluster. Eight of these models have already been published in the literature. We determined a scale height of 1.6 kpc for our exponential thick disk model M1, where we used only sightlines with $|b|$ $>$ 32.5\\degr\\ to avoid the more complex structure at smaller Galactic latitudes. In models M2 and M3 we replaced the thick disk from the models by \\citet{taylor1993} and \\citet{cordes2002} by exponential thick disks. First we had to correctly account for the DM contributions by the other components in these models, which we did by subtracting these DMs from the observed DMs before fitting a new scale height for the thick disk. The thick disks in M2 and M3 have scale heights of 1.6 and 1.3 kpc, resp. When we limit the radial extent of the thick disk in model M2 to 27.7 kpc, it produces for $|b|$=0\\degr\\ the same DM at infinity as the original thick disk from the model by \\citet{taylor1993}. Although this radius on average reproduces the DM of the original thick disk, the DM difference between the old and new thick disk varies somewhat as a function of Galactic longitude. With model M4 we explore using the emission measure as a proxy to predict the observed scatter in DM that the other models cannot explain.  A few of the components in the model by \\citet{cordes2002} contribute also to the DMs of sightlines that lie further than 40\\degr\\ from the Galactic plane. Since \\citet{gaensler2008} only fitted an exponential thick disk, incorporating their model for the thick disk in the model by \\citet{cordes2002} leads to internal inconsistencies.  We developed a novel approach for comparing the predictive quality of the different models, by first calculating for each sightline the factor N by which the modelled DM deviates from the observed DM (so N = modelled DM/observed DM when that ratio is larger than 1, otherwise N = observed DM/modelled DM), and then comparing the cumulative distributions of N between the models. Most models can predict DM to within a factor of 1.5-2 of the measured value for 3/4  lines of sight, and the model by \\citet{taylor1993} that we updated with a more extended thick disk consistently performs better than the other models that we tested. Our fourth model predicts DM least accurately. Surprisingly, the simpler models that we tested perform almost as well as the more complex models. This is probably related to the sparse spatial distribution of the sightlines that we used. Finally, by comparing which pulsar DMs are difficult to predict by most models we identified pulsars that lie beyond \\hii\\ regions. The accurate pulsar distances then place upper limits on the distances to these \\hii\\ regions.  The analysis we present here will benefit from having a larger sample of sightlines with known distances and dispersion measures, which can come from either parallax measurements of additional pulsars, or from discoveries of more globular clusters with pulsars in them. A different approach in the modelling philosophy, for example by including H$\\alpha$ data or a different proxy, warrants further investigation, as it can help explain the detailed structure in DM."
    },
    "1208/1208.1795_arXiv.txt": {
        "abstract": "We propose a new method to conserve the total energy to round-off error in grid-based codes for hydrodynamic simulations with self-gravity. A formula for the energy flux due to the work done by the the self-gravitational force is given, so the change in total energy can be written in conservative form. Numerical experiments with the code Athena show that the total energy is indeed conserved with our new algorithm and the new algorithm is second order accurate. We have performed a set of tests that show the numerical errors in the traditional, non-conservative algorithm can affect the dynamics of the system. The new algorithm only requires one extra solution of the Poisson equation, as compared to the traditional algorithm which includes self-gravity as a source term. If the Poisson solver takes a negligible fraction of the total simulation time, such as when FFTs are used, the new algorithm is almost as efficient as the original method. This new algorithm is useful in Eulerian hydrodynamic simulations with self-gravity, especially when results are sensitive to small energy errors, as for radiation pressure dominated flow. ",
        "introduction": "Self-gravity plays an important role in many astrophysical flows. For instance, it is central to massive star formation \\citep[e.g.,][]{MckeeOstriker2007,Krumholzetal2009}, supernova explosions \\citep[e.g.,][]{Nordhausetal2010}, planet formation in protoplanetary disks \\citep[e.g.,][]{Armitage2011}, and structure formation in the early universe \\citep[e.g.,][]{Bertschinger1998}. Numerical modeling of these systems requires that self-gravity be implemented correctly in hydrodynamical simulations, especially in problems for which energy balance is important. For example, the fate of a collapsing cloud depends on the total energy which is the sum of the potential, internal, and kinetic energies. For some systems, e.g. those having a relativistic equation of state with adiabatic gamma, $\\gamma \\approx 4/3$, the total energy can be much smaller than either the potential or internal energies in hydrostatic equilibrium. Thus, a small numerical error in the computation of the potential energy can lead to a large fractional error in the total energy, potentially causing a bound system to become unbound \\citep[e.g.,][]{JiangGoodman2011} or vice versa. Energy conservation in a grid-based code with self-gravity is usually not guaranteed. The goal of this paper is to propose a new algorithm for grid-based codes, which conserves the sum of the internal, kinetic, and potential energies to round-off error, allowing accurate simulations of self-gravitating systems to be performed.  The change in the momentum and the energy due to self-gravity are typically added to the equations of hydrodynamics as source terms. In this case, the Euler equations read \\begin{eqnarray} \\label{rhoeq} \\frac{\\partial\\rho}{\\partial t}+\\bfnabla\\cdot(\\rho \\bv)&=&0,\\\\ \\label{peq} \\frac{\\partial( \\rho\\bv)}{\\partial t}+\\bfnabla\\cdot({\\rho \\bv\\bv+\\bP}) &=&-\\rho\\bfnabla\\phi, \\ \\\\ \\label{eeq} \\frac{\\partial{E}}{\\partial t}+\\bfnabla\\cdot\\left[(E+P)\\bv)\\right]&=&-\\rho\\bv\\cdot\\bfnabla\\phi, \\ \\\\ \\label{poisson} \\bfnabla^2\\phi=4\\pi G\\rho. \\end{eqnarray} Here $\\rho,\\ \\bv, P$ are density, velocity, and pressure respectively, and $E$ is the sum of the internal and kinetic energies. $G$ is the gravitational constant, and $\\phi$ is the gravitational potential, which is related to the density via the Poisson equation (\\ref{poisson}). We can also define the total energy as \\ba E_{tot} \\equiv E + \\frac{1}{2} \\rho \\phi. \\ea $E_{tot}$ is a {\\it globally} conserved quantity so that \\ba \\label{globalcons} \\frac{\\partial}{\\partial t} \\int d^3{\\bf x} E_{tot} = 0 \\ea for any isolated self-gravitating system. Expression (\\ref{globalcons}) can be derived by integrating equation (\\ref{eeq}) over all of space and applying the divergence theorem (see e.g., \\citealt{BinneyTremaine2008}, section 2.1).  Equations (\\ref{peq}) and (\\ref{eeq}) are written in non-conservative form. Therefore, momentum and energy are usually not conserved to roundoff error when they are solved numerically. It is well known that if momentum is not conserved for a self-gravitating system, then the shape of the object cannot be kept very well when it moves across the simulation box \\citep[e.g.,][]{Edgareal2005,Taskeretal2008}. In order to conserve momentum, it is now common practice to write the momentum equation in conservative form as \\begin{eqnarray} \\frac{\\partial( \\rho\\bv)}{\\partial t}+\\bfnabla\\cdot({\\rho \\bv\\bv+\\bP+\\bT}) =0, \\end{eqnarray} where the gravitational tensor $\\bT$ is \\begin{eqnarray} \\bT=\\frac{1}{4\\pi G}\\left[\\bfnabla\\phi \\bfnabla \\phi-\\frac{1}{2}(\\bfnabla\\phi)\\cdot(\\bfnabla\\phi)\\bI\\right], \\label{equation_momentum} \\end{eqnarray} and $\\bI$ is the identity tensor. This equation is mathematically equivalent to the original momentum equation (\\ref{peq}) with use of equation (\\ref{poisson}). However equation (\\ref{equation_momentum}) can be solved in a completely different way numerically. Momentum change due to gravitational acceleration is no longer included as a source term. Instead, it is included as a flux term via $\\bT$, which means the total momentum can be conserved to round-off error numerically. This has already been implemented in many grid-based codes, such as Athena.  The improvement we propose in this paper is to also write the energy equation with self-gravity (\\ref{eeq}) in conservative form. Using this form of the equation, we devise a numerical algorithm that conserves the total energy with self-gravity to round-off error. ",
        "conclusions": "We have developed a new algorithm to conserve total energy to round-off error for Eulerian hydrodynamical simulations incorporating self-gravity. We have implemented this new algorithm in Athena and shown that it conserves total energy to round-off error as expected. By comparing a set of tests in 1D, 2D and 3D with two different numerical algorithms for self-gravity (the new conservative algorithm and the traditional non-conservative algorithm), we have shown that the numerical error made in the traditional algorithm can change the dynamics significantly. From our numerical experiments, we conclude that the conservative algorithm will be important when a small amount of energy error can dramatically affect the dynamics, such as in radiation-dominated systems with an adiabatic index close to $4/3$.  We tested our new algorithm by implementing it in Athena, which uses an unsplit Godunov method to solve the equations of hydrodynamics. In principle, this algorithm can also be used in codes that implement an operator split scheme, such as ZEUS \\citep[][]{StoneNorman1992}. In the application using ZEUS with self-gravity performed by \\cite{JiangGoodman2011} , the energy error using the traditional non-conservative algorithm was positive. If ZEUS is used to study problems of star formation with the non-conservative algorithm, gravitationally bound clumps can acquire a positive total energy due to numerical errors and be destroyed as a result of numerical heating. However, if our conservative algorithm is used, this will not happen, since the total energy is conserved.  The new conservative algorithm is not necessary for all hydrodynamical simulations incorporating self-gravity. We find that for cases when the energy error is not a concern, the two algorithms give almost identical results. For example, for small oscillations of a polytropic sphere with an adiabatic index of $\\gamma=5/3$, the two algorithms yield very similar temporal and spatial behavior. In such cases, the old algorithm is just as good as the new one and is potentially more efficient."
    },
    "1208/1208.3103_arXiv.txt": {
        "abstract": "Based on an analysis of data obtained with the Wide Field Camera 3 (WFC3) on the Hubble Space Telescope (HST) we report the identification of two distinct stellar populations in the core of the giant H{\\sc ii} region 30~Doradus in the Large Magellanic Cloud. The most compact and richest component coincides with the center of R136 and is $\\sim 1$ Myr younger than a second more diffuse clump, located $\\sim 5.4$ pc toward the northeast. We note that published spectral types of massive stars in these two clumps lend support to the proposed age difference. The morphology and age difference between the two sub-clusters suggests that an ongoing merger may occurring within the core of 30~Doradus. This finding is consistent with the predictions of models of hierarchical fragmentation of turbulent giant molecular clouds, according to which star clusters would be the final products of merging smaller sub-structures. ",
        "introduction": "As  the most luminous and best known massive star-forming region in the Local Group \\citep[LG,][]{Kennicutt1984}, the Tarantula Nebula in the Large Magellanic Cloud (LMC) offers us the unique opportunity to investigate the process of star formation (SF) in an environment that, in many ways (such as metallicity, dust content, and star formation rate), resembles the extreme conditions of the early universe and distant star-forming regions. The star formation history (SFH) of Tarantula's ionizing cluster NGC~2070 (a.k.a. 30~Dor) is complex. \\citet{Walborn1997} defined 30~Dor as a two-stage starburst, in which the radiation from the compact core R136 has triggered a new generation of stars at a projected distance of $\\sim 1\\arcmin$ (corresponding to $\\sim 15\\, {\\rm pc}$ at the the distance of the LMC).  The complexity of NGC~2070 is reflected also in its atypical luminosity profile, which shows a break at  $\\sim 10\\arcsec$ ($\\sim 2.5\\, {\\rm pc}$) and a bump at $\\sim 30\\arcsec$ \\citep[$\\sim 7\\, {\\rm pc}$;][]{Mackey2003}. More recent observations in the near IR with Multiconjugate Adaptive Optics (MAD) indicate that the innermost $\\sim30\\arcsec$ of the luminosity profile can be fitted with a single \\citet[][EFF]{Elson1987} powerlaw, once the asymmetric shape of 30~Dor is taken into account \\citep{Campbell2010}.  Here we take advantage of the high spatial resolution and sensitivity of the Wide Field Camera 3 (WFC3) on board of the {\\it Hubble} Space Telescope (HST) to study the innermost $\\sim2\\arcmin 44\\arcsec\\times 2\\arcmin 44\\arcsec$ of NGC~2070. In particular, we use public UVIS and IR channel observations (P.I. R.W. O'Connell, GO-11360) to characterize the spatial distribution of its stellar populations. ",
        "conclusions": "We used deep HST/WFC3 UV, optical and near-IR observations of the core of 30~Dor to study its stellar content. In all the filters we identified a dual structure that cannot be attributed to dust extinction. A careful analysis of the stellar content of the clumps indicates that, while the Northeast clump formed the majority of its stars between 2 and 5 Myr ago, R136 started to form stars likely $\\sim 2\\, {\\rm Myr}$ ago and was still active $\\sim 1\\, {\\rm Myr}$ ago.  In retrospect, there are substantial indications in the literature that the Northeast Clump comprises a distinct, more evolved entity.  Its two brightest members, R141 and R142, are early B supergiants that must be older than the denizens of R136 itself \\citep{Feast1960, Parker1993}. Many of the fainter late-O/early-B stars, being evolved,  are also subject to the same age argument \\citep{Walborn1997, Massey1998, Bosch1999}. In an analysis of star formation in the 30~Dor core based upon these observations, \\citet{Selman1999} proposed three sequential events with mean ages less than 1.5~Myr, 2.5~Myr, and 5~Myr; their Figure~9 shows that the first is strongly concentrated toward R136, while the heaviest concentration of the intermediate and old objects coincides with the Northeast Clump.  The age difference between R136 and the Northeast Clump can also explain  the, until now, puzzling presence of two WC stars in 30~Dor, as these highly evolved objects are seldom found in giant H{\\sc ii} regions.  One of them, Mk~33d, is located in an apparent multiple system at the inner edge of the Northeast clump \\citep{Melnick1985}, while the northward tail on the Northeast clump encompasses the multiple system R140 containing the other WC along with two WN's \\citep{Moffat1987}.  Finally, even the isolated M supergiant  IR18 \\citep{McGregor1981} somewhat further to the NE, long considered a field interloper, could actually be associated with the older clump. Positions of the mentioned stars with respect to the two sub-clusters are shown in Fig.~\\ref{f:zoom}.  This is, of course,  a complex three-dimensional region, and some earlier spectral types also appear in the predominantly older fields, particularly in the Mk~33 system.  The O2~If*/WN5 object Mk~30 \\citep{Crowther2011} that lies adjacent to R142 along the ridge of the Northeast clump could be associated with R136; but on the other hand, some such objects are also found in older young clusters \\citep[e.g., TS3 in NGC~2060;][]{Schild1992}, possibly as a result of binary evolution. In summary, the information presented here provides strong support for our completely independent recognition of a double system with an age offset between them in the core of 30~Dor.  Characterizing the complex structure of massive star-forming regions such as NGC~2070 is important to understand the early stages of cluster formation. The age difference found between R136 and the Northeast clump, as well as their morphology, may indicate that the core of NGC~2070 is the result of a recent or ongoing merger between two sub-clusters. Interestingly from the analysis of the radial velocities of apparently single O stars, \\citet[][submitted]{Brunet2012} found evidence for an internal rotation in NGC~2070, that can be due to either a recent merger between the main core of NGC~2070 and a secondary cluster or by the agglomeration and clumpiness of the gas.  The majority of the stars studied in \\citet{Brunet2012} coincides with R136, but 5 sources coincide with the Northeast Clump.  If the total mass of the system is $\\sim 10^5\\, {\\rm M_\\odot}$ \\citep{Andersen2009} then the velocity difference, assuming a distance $d=4.5\\, {\\rm pc}$, would be $v=\\sqrt(2GM/d)\\simeq 13\\, {\\rm Km s^{-1}}$ for a hyperbolic encounter (i.e. 0 energy), and $\\simeq 9\\, {\\rm km s^{-1}}$ for a circular orbit. The maximum velocity difference reported by \\citet{Brunet2012} is $\\sim 5\\,  {\\rm km s^{-1}}$. If R136 and the Northeast Clump are undergoing a merger, then either we are observing the interaction  with a low inclination (which would agree with the fact that we see the tidal tail on the plane of the sky), or the 2 systems are bound and are falling in with low velocity. A detailed analysis of the mass distribution in the two sub-clusters and a more sophisticated modeling of the two component interaction will be presented in two forthcoming papers (Sabbi et al. in prep, Gieles et al. in prep.)  The presence of distinct components in a very young system is also consistent with recent studies of galactic giant molecular clouds (GMC) \\citep{Heyer2009}, which contradict the traditional paradigm that GMCs are gravitationally bound systems. \\citet{Ballesteros2011} suggested that GMCs can undergo a hierarchical gravitational collapse, with the collapse occurring on scales ranging from individual cores to the whole cloud. In this scenario, because of the turbulent fragmentation of a GMC, SF will not be spread over the whole GMC, but will be localized in gravitationally bound pockets of gas \\citep{Clark2005, Clark2008}. Star clusters will be the final products of the hierarchical merger of such smaller sub-structures \\citep{Bonnell2003, Bate2009, Federrath2010}.  We finally note that in a recent paper,  \\citet{Fujii2012} suggested that the cuspy density profile of R136 and the large number of massive runaway stars that seem to escape from this cluster are signatures of a post-core-collapse star cluster \\citep{Mackey2003}. Considering the small age of R136, such an early core collapse would occur only  if the cluster had an initial density of $\\rho_c\\ge 10^6 {\\rm M_\\odot}\\, {\\rm pc^{-3}}$, which is considerably higher than the current estimates.  \\citet{Fujii2012}  suggest that an efficient way to speed up the collapse process is to invoke the hierarchical merging of several smaller clusters.  Observational evidence for a hierarchical process of SF comes also from the complex structure of massive star-forming regions such as NGC~604 in M33 \\citep{jesus}, or NGC~346 in the Small Magellanic Cloud \\citep{sabbi08}. Also, \\citet{Gennaro2011} suggested that the elongation and mass segregation found in Westerlund~1 may be the footprint of a merger of multiple sub-clusters, formed almost coevally in the parental GMC. If massive star clusters form from mergers of smaller sub-systems, this can also explain the high fraction of rotating globular clusters found in our Galaxy. A study of the radial velocities and/or proper motions of the stars in the two clumps would provide a further test of the merging scenario. Such datasets are currently being obtained by us (e.g. H\\'enault-Brunet et al. submitted; Lennon et al. in prep.)."
    },
    "1208/1208.4723_arXiv.txt": {
        "abstract": "{ We present results of spectral investigations of the Sco\\th X-1 like Z-track sources Sco\\th X-1, GX\\th 349+2 and GX\\th 17+2 based on {\\it Rossi-XTE} observations determining spectral evolution along the Z-track. Results are obtained for a spectral model describing an extended accretion disk corona. The results are compared with previous results for the Cyg\\th X-2 like group: Cyg\\th X-2, GX\\th 340+0 and GX\\th 5-1 and a general model for the Z-track sources proposed. On the normal branch, the Sco-like and Cyg-like sources are similar, the results indicating an increase of mass accretion rate $\\dot M$ between soft and hard apex, not as in the standard view that this increases monotonically around the Z. In the Cyg-like sources, increasing $\\dot M$ causes the neutron star temperature $kT$ to increase from $\\sim$1 to $\\sim$2 keV. At the lower temperature, the radiation pressure is small, but at the higher $kT$, the emitted flux of the neutron star is several times super-Eddington and the high radiation pressure disrupts the inner disk and launches the relativistic jets observed on the upper normal and horizontal branches. In the Sco-like sources, the main physical difference is the high $kT$ of more than 2 keV on all parts of the Z-track suggesting jets are always possible, even on the flaring branch. The flaring branch in the Cyg-like sources is associated with a release of energy on the neutron star consistent with unstable nuclear burning. The Sco-like sources are very different as flaring appears to be a {\\it combination} of unstable burning and an increase of $\\dot M$ which makes flaring much stronger. Analysis of 15 years of {\\it RXTE} ASM data on all 6 classic Z-track sources shows the high rate and strength of flaring in the Sco-like sources suggesting that the continual release of energy heats the neutron star causing the high $kT$. Analysis of a Sco\\th X-1 observation with unusually little flaring supports this. GX\\th 17+2 appears to be transitional between the Cyg and Sco-like types. Our results do not support the suggestion that Cyg or Sco-like nature is determined by the luminosity. ",
        "introduction": "The Z-track sources form the brightest group of Low Mass X-ray Binaries (LMXB) containing a neutron star, with luminosities at or above the Eddington limit. The sources display three tracks in hardness-intensity diagrams: the Horizontal Branch (HB), Normal Branch (NB) and Flaring Branch (FB) (Hasinger \\& van der Klis 1989) showing that major physical changes take place within the sources. However, the nature of these changes has not been understood and is essential to an understanding of LMXB in general. It has been widely thought that the physical changes are caused by change of the mass accretion rate along the Z-track (Priedhorsky et al. 1986), assumed to increase in the direction HB$\\rightarrow$NB$\\rightarrow$FB, however, the evidence for this is rather limited (below). Moreover, the variation of X-ray intensity does not suggest this, increasing on the HB, but then decreasing on the NB.  It was apparent that the sources fell into two groups (Kuulkers et al. 1997; Hasinger \\& van der Klis 1989): firstly, the Cyg-like sources Cygnus\\th X-2, GX\\th 5-1 and GX\\th 340+0 displaying full Z-tracks with HB, NB and FB but having weak flaring. Secondly in the Sco\\th X-1 like group: Sco\\th X-1, GX\\th 349+2 and GX\\th 17+2 there is hardly any HB but flaring is strong and frequent with large increases of X-ray intensity. These substantial differences have not been understood.  The sources are important because of the observation of relativistic jets in the radio, notably the charting of outward motion of material from Sco\\th X-1 by Fomalont et al. (2001) at a velocity $v/c$ $\\sim$0.45. As jets are seen predominantly in the HB (e.g. Penninx 1989), X-ray observations allow us to compare conditions at the inner disk when jets are formed and are not formed, allowing us to investigate jet launching observationally as opposed to theoretically, aiming to understand the disk-jet connection.  Since the discovery of the sources as a class with three basic states (Schulz et al. 1989) there has been relatively little progress in understanding them. Extensive timing studies have been carried out (van der Klis et al. 1987; Hasinger et al. 1989) revealing complex behaviour of quasi-periodic oscillations (QPO) at low frequencies ($\\sim$6 Hz) and at kilohertz frequencies; see review of van der Klis (2003). Spectral investigations may be more likely to provide an explanation, as revealing physical changes in the emission components during motion along the Z-track, but have been hindered by lack of agreement on spectral models.  \\subsection{The extended ADC}  The broadband spectra exhibit both non-thermal and thermal emission, which have been interpreted using two very different physical models as have the spectra of LMXB in general. In the Eastern model (Mitsuda et al. 1989), non-thermal emission originates in a small central region and thermal emission from the inner disk. It has been widely applied to the Z-track sources, however, there is no generally accepted physical interpretation of the resulting parameter variations.  The Extended Accretion Disk Corona model (Church \\& Ba\\l uci\\'nska-Church 1995, 2004) was developed out of the original Western model in which blackbody emission is from the neutron star and non-thermal emission is Comptonization from an ADC above the disk. The extended nature of the accretion disk corona (ADC) was revealed by work on the dipping class of LMXB, sources having high inclination such that decreases of X-ray intensity are seen in every binary orbit, due to absorption in the bulge in the outer disk where the accretion flow impacts (White \\& Swank 1982; Church et al. 1997).  The dipping sources strongly constrain emission models since spectra of non-dip and several levels of dipping must all be fitted. The gradual removal in dips of the dominant Comptonized emission over a time of order 100 seconds argues strongly for the emission being extended. Measurements of the dip ingress time in the dipping sources show that the ADC has radial sizes $R$ of 10\\th 000 to 700\\th 000 km depending on luminosity, and is thin, i.e with a height $H$ having $H/R$ $<<$ 1: a thin corona above a thin disk (Church \\& Ba\\l uci\\'nska-Church 2004). The spectral evolution in dipping shows that the Comptonized emission is removed gradually by an extended absorber having increasing overlap with the extended source, again demonstrating the extended nature of the ADC. The Extended ADC model gave very good fits to the complex spectral evolution in dipping in many observations of the dipping LMXB (Church et al. 1997; 1998a,b, 2005; Ba\\l uci\\'nska-Church et al. 1999, 2000; Smale et al. 2001; Barnard et al. 2001). In addition it was able to fit the spectra of all types of LMXB in a survey made with {\\it ASCA} (Church \\& Ba\\l uci\\'nska-Church 2001) and so capable of describing sources of all inclinations.  Recently support for the extended ADC came from the {\\it Chandra} grating results of Schulz et al. (2009) by an independent technique. Precise grating measurements revealed a wealth of emission lines of highly excited states (many H-like ions) such as Ne X, Mg XII, Si XIV, S XIV, S XVI, Fe XXV and Fe XXVI. The line widths seen as Doppler shifts due to orbital motion indicated emission in a hot ADC at radial positions between 18\\th 000 and 110\\th 000 km in good agreement with the overall ADC size from dip ingress timing.  In multi-wavelength observations of Cygnus\\th X-2 in 2009 (Ba\\l uci\\'nska-Church et al. 2011) a series of absorption edges below 1.4 keV remarkably had no detectable increase in depth during dips, showing that part of the continuum emission had to be extended so that part was not covered, while covered emission below 1.4 keV was completely removed.  Thus the evidence favours an extended ADC and this model has provided an explanation of the Cygnus\\th X-2 sub-group (below) which seems convincing and has the potential to lead to a better understanding of the sources and jet formation. In the present work we test the hypothesis that the model can provide a physically reasonable explanation of the Sco\\th X-1 like sub-group.  \\subsection{A model for the Cygnus\\th X-2 like Z-track sources}  We previously applied the Extended ADC model to high-quality {\\it Rossi-XTE} data on the Cyg-like sources GX\\th 340+0 (Church et al. 2006; Paper I), GX\\th 5-1 (Jackson et al. 2009; Paper II) and Cygnus\\th X-2 (Ba\\l uci\\'nska-Church et al. 2010; Paper III). In all cases the model fitted the spectra well at all positions on the Z-track, and clearly suggested a physical explanation for the Z-track phenomenon. All three sources behaved in the same way showing the explanation to be consistent. In this model, the Soft Apex between the NB and FB has a low mass accretion rate $\\dot M$ and so a low X-ray intensity. Moving along the NB towards the Hard Apex, there is a strong increase in luminosity, 90\\% of which is due to the ADC Comptonized emission and it is difficult to see how this can happen without an increase of mass accretion rate. The resultant increase in $kT_{\\rm BB}$ from $\\sim$1 keV at the soft apex to $\\sim$2 keV at the hard apex and higher values still on the HB results in a large increase in $T_{\\rm BB}^4$ and so a large increase in the radiation pressure of the neutron star. The ratio $f/f_{\\rm Edd}$: the flux emitted by unit area of the star $f$ as a fraction of the Eddington value ($L_{\\rm Edd}/4\\,\\pi\\,R_{\\rm NS}^2$, where $R_{\\rm NS}$ is the radius of the neutron star) increases from $\\sim$0.3 at the soft apex to 3 times super-Eddington on the HB. The good correlation of high radiation pressure with the radio emission of jets on this part of the Z-track led us to suggest that high radiation pressure disrupts the inner disk deflecting the accretion flow vertically so launching the jets.  On the HB $L_{\\rm ADC}$ decreases suggesting $\\dot M$ decreases while $kT_{\\rm BB}$ continues to rise. It is not definite what happens: $\\dot M$ may be decreasing in the disk affecting $L_{\\rm ADC}$ but not yet reaching the neutron star; alternatively, the high radiation pressure may be disrupting the disk out to larger radial positions so reducing the Comptonized emission.  In flaring, $L_{\\rm ADC}$ was constant, suggesting that $\\dot M$ was constant, but there is a strong increase in $L_{\\rm BB}$. If $\\dot M$ does not change, flaring must be due to an additional energy source on the neutron star, i.e. unstable nuclear burning. Support for this came from measured values of the mass accretion rate per unit emitting area on the neutron star $\\dot m$, the critical parameter in the theory of stable/unstable burning on the surface of the neutron star (e.g. Bildsten 1998). The theory describes various r\\'egimes of stable and unstable nuclear burning over the wide range of luminosities found in LMXB, including unstable burning at lower luminosities seen as X-ray bursts but is also applicable at the luminosities of the Z-sources. In all three sources, $\\dot m$ at the onset of flaring (the soft apex) agreed with the theoretical critical value below which there is unstable helium burning.   \\subsection{A monotonic variation of mass accretion rate ?}  This model for the Cygnus\\th X-2 like sources with $\\dot M$ increasing on the NB and constant on the FB does not support the standard view that $\\dot M$ increases monotonically HB$\\rightarrow$NB$\\rightarrow$FB. In a multi-wavelength campaign on Cyg\\th X-2 (Hasinger et al. 1990),  an increase of UV emission (Vrtilek et al. 1990) in what appeared to be flaring suggested an $\\dot M$ increase. However, there has been some confusion in the Cyg-like sources between dipping and flaring in that {\\it reductions} in X-ray intensity in dipping (seen in hardness-intensity) when plotted in colour-colour look like a FB. In Cyg\\th X-2 recent multi-wavelength observations show that intensity reductions were indeed absorption dips unrelated to flaring (Ba\\l uci\\'nska-Church et al. 2011, 2012). In GX\\th 340+0 and \\hbox{GX\\th 5-1}, a 4th track has been seen at the end of the FB with decreasing intensity suggesting a dipping-flaring link. However, spectral fitting has shown that this was not absorption but a reduction in the Comptonized emission level (Church et al. 2010). Thus the original evidence for the monotonic increase of $\\dot M$ becomes questionable.   Support for the monotonic increase of $\\dot M$ was found in the timing properties of the sources with many properties correlating with Z-track position, i.e. arc length $S$ along the track. However, in recent years authors have found several properties that do not correlate (Wijnands et al. 1996; Homan et al. 2002; Casella et al. 2000; Dieters \\& van der Klis 2000) so that $S$ is not a measure of $\\dot M$.   In our model for the Cyg-like sources, it is proposed that $\\dot M$ increases between the soft and hard apex. In the present work  we will test this proposal in the case of the Sco\\th X-1 like sources, and find that this is also the case.  \\subsection{Approach of the present work}  In this work, we test the hypothesis that the Extended ADC model can describe the Sco\\th X-1 like sources, provide a physically reasonable explanation of them and of the differences between these and the Cyg-like sources.  It is applied in the form {\\sc bb + cpl} where {\\sc bb} is simple blackbody emission of the neutron star and {\\sc cpl} is a cut-off power law representing Comptonization in an extended corona. As discussed in detail in Sect. 3.3, the model was applied with a fixed power law photon index of 1.7 and a free column density (except for Sco\\th X-1). The {\\sc bb + cpl} spectral form is appropriate to Comptonization of seed photons from the disk below the ADC which at large radial positions have low energies ($kT$ $\\sim$ 0.1 keV or less); see Church \\& Ba\\l uci\\'nska-Church (2004) for a full discussion. A simple blackbody model is used and we have not previously found any evidence for departure from this form. Theoretical treatments of electron scattering in the neutron star atmosphere require measured blackbody temperatures to be corrected. Observational evidence for a modified blackbody is somewhat limited as discussed in Ba\\l uci\\'nska-Church et al. (2001) who show that whether modification is important depends critically on the electron density which is poorly known. However, we show in Sect. 3.4 that this correction combines with the relativistic correction to have a relatively small effect on the blackbody temperature and radius.   \\subsection{The sources: Sco\\th X-1, GX\\th 349+2 and GX\\th 17+2}  Sco X-1 is the brightest persistent X-ray source in the sky (Giacconi et al. 1962) at a distance of 2.8$\\pm $0.3 kpc (Bradshaw et al. 1999) and has been studied intensively since the early 1960s, including multi-wavelength campaigns (Canizares et al. 1975; Hertz et al. 1992). It has an orbital period of 18.9 hr (Cowley \\& Crampton 1975) and a M type companion of mass 0.4 $M_{\\sun}$ (Steeghs \\& Casares 2002). Spectral modelling of Sco\\th X-1 (e.g. White et al. 1985) showed that the spectra could be fitted by a blackbody plus Comptonization model. The third branch, the horizontal branch, was discovered by Hasinger et al. (1989) using {\\it Exosat}. The first kHz QPO was discovered in this source by van der Klis et al. (1996).  GX\\th 349+2 is relatively little studied but has a similar period of 21.85$\\pm 0.4$ hr (Wachter \\& Margon 1996). It has not been seen in the HB. A distance of 9.2 kpc was found by Penninx (1989) and van Paradijs \\& White (1995) (see also Grimm et al. 2002). Zhang et al. (1998) discovered kHz QPO on the upper normal branch at 712 and 978 Hz. Iaria et al. (2004) found emission features probably L-shell emission of Fe XXIV, Ly$\\,{\\alpha}$ S XVI and Fe XXV.  GX\\th 17+2 appeared as a source with two branches (the NB and HB) in hardness-intensity (Schulz et al. 1989). A strong FB was seen by Hasinger \\& van der Klis (1989). The companion was difficult to identify (Deutsch et al. 1999; Callanan et al. 2002). Investigations of the variation around the Z-track of low frequency QPO were carried out by Langmeier et al. (1990), Kuulkers et al. (1997) and Homan et al. (2002). Kilohertz QPO were discovered by Wijnands et al. (1997). It is one of the two Z-track sources with occasional X-ray bursts: GX\\th 17+2 and Cyg\\th X-2 (Kahn \\& Grindlay 1984; Tawara et al. 1984; Sztajno et al. 1986; Kuulkers et al. 1997, 2002). Migliari et al. (2007) from simultaneous radio and X-ray observations show the switch-on of the jet on the normal branch. Distances of about 7.5 kpc were found by several authors (e.g. Ebisuzaki et al. 1984); distances used in the analysis following are given in Table 6 and further discussed in Sect. 4.5. ",
        "conclusions": "\\subsection{Differences between the Sco\\th X-1 like and the Cyg\\th X-2 like sources}  \\begin{figure}[!h]                                                    % \\begin{center} \\includegraphics[width=84mm,height=84mm,angle=270]{fig12}        % \\caption{The major difference between the Sco\\th X-1 like and Cyg\\th X-2 like sources at the soft apex; upper panel: the blackbody temperature $kT_{\\rm BB}$; lower panel: the corresponding blackbody radii $R_{\\rm BB}$. The dotted lines correspond to the upper and lower limits of the average $R_{\\rm BB}$ for the three Cyg-like sources of 11.6$\\pm$0.6 km at 90\\% confidence.} \\label{} \\end{center} \\end{figure}  On the normal branch, the Sco-like and Cyg-like sources are similar. In the latter, we suggested that the large increase of $L_{\\rm ADC}$ was due to increasing mass accretion rate. $L_{\\rm ADC}$ also increases in the Sco-like sources on the NB (Fig. 7) supporting the idea that $\\dot M$ increases on the NB in all Z-track sources, unlike the often-held standard view that $\\dot M$ increases monotonically in the opposite direction HB$\\rightarrow$NB$\\rightarrow$FB.  The neutron star temperature in the Cyg-like sources is lowest at the soft apex and the blackbody radius maximum (for the non-flaring source) suggesting that the whole star is emitting. The average value for the 3 sources is 11.6$\\pm$0.6 km, but should not be taken literally as the neutron star radius. Firstly, a relativistic correction (1 + z) of 34\\% would reduce the value to 8.7 km. Also applying a colour correction (Sect. 3.4) would however, increase the final value to 18.7$\\pm$0.97 km which is disturbingly high. In addition there may be possible systematic effects such as blocking of part of the neutron star emission by the disk or corona, although this is not generally seen in LMXB except in dipping. Cygnus\\th X-2 is the only source with dipping and high inclination; however, $R_{\\rm BB}$ at the soft apex in this source does not appear different from the other Cyg-like sources.    It was proposed that increasing radiation pressure on the NB disrupted the inner disk leading to accretion on an equatorial belt, and possible mechanisms for this suggested (Paper I).  In Sco\\th X-1 and GX\\th 349+2, $kT_{\\rm BB}$ is substantially higher, this constituting the first major difference, and $R_{\\rm BB}$ is reduced. To emphasize the differences in Fig. 12 we show $kT_{\\rm BB}$ and $R_{\\rm BB}$ for all 6 sources at the soft apex. The reduction of $R_{\\rm BB}$ in the Sco-like sources at the soft apex with higher $kT_{\\rm BB}$ is consistent with the reduction in the Cyg-like sources with increasing $kT_{\\rm BB}$ on the NB.  Closely related is the radiation pressure of the neutron star emission expressed via $f/f_{\\rm Edd}$ using the measured flux (Fig. 8). In Sco\\th X-1 and GX\\th 349+2 the values are always super-Eddington whereas the Cyg-like sources have sub-Eddington values at the soft apex becoming high at the hard apex where jets are observed. Thus in \\hbox{Sco\\th X-1} and GX\\th 349+2 in which the radiation pressure is always high we expect the sources to be permanently subject to disruption of the inner disk and jet launching may take place even from the soft apex although $f/f_{\\rm Edd}$ is only unity. Whether this occurs requires simultaneous radio and X-ray observations during Z-track movement. We might also possibly expect jets during flaring on the basis that $f/f_{\\rm Edd}$ remains high.  GX\\th 17+2 appears transitional not having a high $kT_{\\rm BB}$ at the soft apex nor a very reduced  $R_{\\rm BB}$. Flaring is however, similar to that in Sco\\th X-1 and GX\\th 349+2 with a definite increase of Comptonized emission.  The second major difference is in flaring. The Cyg-like sources have constant $L_{\\rm ADC}$ arguing for a constant $\\dot M$ but an increase of blackbody luminosity. In the Sco-like sources, $L_{\\rm ADC}$ increases and the FB appears more similar to the NB (Fig. 7). In the Cyg-like sources, we proposed an additional energy source of unstable nuclear burning on the neutron star giving increased $L_{\\rm BB}$ at constant mass accretion rate. This was supported by good agreement of the mass accretion rate per unit area $\\dot m$ at the soft apex with the theoretical $\\dot m_{\\rm ST}$ at the border of unstable/stable He burning in bright LMXB in a mixed H/He environment (e.g. Bildsten et al. 1998). For a source descending the NB, at the soft apex conditions in the neutron star atmosphere allow unstable burning and a flare begins. The Sco-like sources are clearly different with increases in both $L_{\\rm ADC}$ and $L_{\\rm BB}$. This suggests an increase of $\\dot M$ on both NB and FB, while some other factor must be different in flaring and it is natural to think that this may be unstable burning. The increase of blackbody luminosity in flaring in the Sco-like sources is similar to that in GX\\th 340+0 (Fig. 7) supporting this. In all Z-track sources there is an increase of blackbody radius in flaring (Fig. 5) suggestive of spreading of the burning region as might be expected across the neutron star in the Sco-like sources and possibly in an expanding atmosphere in the Cyg-like sources.  In GX\\th 17+2, $\\dot m$ at the soft apex is close to $\\dot m_{\\rm ST}$ (Fig. 9) so that unstable burning is expected in flaring. However, in Sco\\th X-1 and GX\\th 349+2 $\\dot m$ $>$ $\\dot m_{\\rm ST}$ implying that burning will be stable. A possible explanation is that with high radiation pressure at all times in these sources, the mass flow to the neutron star is modified by disruption of the inner disk, with mass permanently diverted out of the disk. $\\dot m$ is obtained from $\\dot M$ assuming the standard relation with total luminosity, so that if mass is diverted away from the neutron star we over-estimate $\\dot M$ but $\\sim$80\\% of the flow would have to be diverted to bring $\\dot m$ into the unstable region.  \\subsection{A model for the Sco\\th X-1 like sources}  The results suggest a model for the Sco-like sources as follows. Ascending the NB the sources are similar to the Cyg-like sources with $\\dot M$ increasing. Flaring is very different as $\\dot M$ also increases suggesting that flaring is unstable nuclear burning {\\it combined} with increasing $\\dot M$. The additional burning process causes the FB in hardness-intensity to deviate from the NB (becoming less hard). Variations of the mass accretion rate sometimes cause the source to move on the NB, but if unstable burning is triggered then the movement is on the FB. Unstable burning might be triggered if part of the mass flow has been diverted into the jet and away from the neutron star so that $\\dot m$ $<$ $\\dot m_{\\rm ST}$. The radiation pressure is higher on the NB (Fig. 8) in the Sco-like sources so a source descending the NB may fulfil the unstable burning condition. Another possibility is that if only part of the neutron star is emitting at the soft apex, on the edge of the emitting region $\\dot m$ may be reduced and unstable burning possible. Galloway (2008) made a similar suggestion that accretion over a limited part of the star might at lower mass accretion rates produce X-ray bursts confined to part of the star, which might spread across the star (Bildsten 1995). The Cyg-like sources at the soft apex accrete over the whole star apparently so unstable burning only occurs when $\\dot m$ falls below the critical value. The Sco-like sources would be expected to move onto the NB from the soft apex after a period of flaring as further accumulation of matter would be needed for the next flare.   Secondly, the high temperature of the neutron star in Sco\\th X-1 and GX\\th 349+2 at all times $>$ 2 keV results from the continual strong flaring releasing energy on the neutron star in contrast with the Cyg-like sources where flaring is rare and weak. Between flares, the neutron star could cool in principle by the emission of X-rays; however, it appears observationally that cooling is not effective on the timescale between flares as discussed below.  Thirdly, the model predicts that jets are formed when the radiation pressure is high. In the Cyg-like sources, this is in the region of the hard apex but not at the soft apex. However, in Sco\\th X-1 the radiation pressure is high even at the soft apex and in flaring. The model thus predicts that radio emission in Cyg\\th X-2 will sometimes be zero, i.e. quenched in flaring but emission in Sco\\th X-1 may not become zero at the soft apex. In Cyg\\th X-2 radio fluxes up to 5 mJy were found on the horizontal and upper normal branch at 1.49, 4.9 and 8.4 GHz but the source was not detected on the lower NB and FB with flux  $<$ 1 mJy, i.e. showing a clear correlation with Z-track position (Hjellming et al. 1990a). Similarly, Ba\\l uci\\'nska-Church et al. (2011) found a 5$\\sigma$ upper limit of $<$ 150 $\\mu$Jy at 5 GHz in e-EVN observations in 2009, May. Sco X-1 on the other hand has both radio-loud and radio-weak states and weak radio emission has been detected when on the FB (Hjellming et al. 1990b). This tends to support the prediction but clearly further testing is needed.   \\subsection{Cooling rate}  We propose that unstable nuclear burning deposits energy in the neutron star such that, between flares, i.e. at the soft apex, the blackbody temperature measured for the atmosphere does not decrease but remains at 2 keV as in a flare. This may be compared with X-ray bursting in which the emission cools rapidly from $\\sim$ 2.2 keV to 1 keV between bursts. It thus appears that radiative losses are not effective in cooling the neutron star between flares. However, we found some evidence for a 10\\% cooling after 3 days without flares. This raises the interesting question of the rate of cooling of the neutron star.  In the case of X-ray burst sources, a small fraction of the nuclear energy released will flow into the interior at each burst, heating it. The time-averaged inward flow will be small: $\\sim$ 4\\% (Hanawa \\& Fujimoto 1984), because most of the energy released is radiated. In equilibrium, there will be an outward energy flow balancing this (Lamb \\& Lamb 1978; Ayasli \\& Joss 1982; Hanawa \\& Fujimoto 1984). Compressional heating and neutrino cooling will also be involved. The averaged inward flow is small, $\\approx 10^{35}$ erg s$^{-1}$, and the thermal capacity of the neutron star dominated by partially degenerate electrons is large, assuming the neutrons are super-fluid (Tsuruta et al. 1972) $\\approx$ $3\\times 10^{44}$ erg (Lamb \\& Lamb 1978). Thus the heating timescale can be $\\sim$100 years.   Unstable burning in flaring will also heat the star but whereas a burst lasts $\\sim$100 seconds and occupies only 1\\% of the interval between bursts, flaring is more prolonged at $\\sim$5000 seconds and the ratio of flare duration to flare spacing is higher. Thus the inward heat flow will be higher and so the heating timescale reduced, and in equilibrium, the outward flow of energy from the interior would also be increased.  To compare with our results, we consider what happens if the flaring is suddenly turned off after a long period of continuous flaring. If we assume an outward flow of heat much higher than in the burst source case, say $10^{37}$ erg s$^{-1}$, for the whole neutron star to cool (i.e. the partially degenerate electrons) by 10\\% in temperature would require $\\approx$0.1 year, about 10 times longer than the probable 3 days timescale. This may suggest that the appropriate cooling rate is not the equilibrium rate of outward heat flow, but a higher rate, for example, electrons in the crust may cool more rapidly than that.             \\subsection{The energetics of flaring}  It may be thought that since nuclear burning, e.g. of hydrogen, is 30 times less efficient than accretion power, flaring cannot be nuclear in nature, and in the standard view of Z-track sources it is increasing $\\dot M$ that powers flares in both Cyg and Sco-like sources. However, in the Cyg-like sources an increase of neutron star luminosity in flaring of $10\\times 10^{37}$ erg s$^{-1}$ for 2000 seconds say, releases a flare energy $E_{\\rm flare}$ of $\\sim 1\\times 10^{41}$ erg. Instantaneous nuclear burning of the accretion flow cannot supply this, but matter accumulating over a period of $\\approx$ 3 hours provides an energy $\\epsilon\\,M\\,c^2$ $\\geq$ $E_{\\rm flare}$. The flares in the Sco-like sources are stronger but only $\\sim$50\\% or less of the energy appears to be nuclear burning. $L_{\\rm BB}$ increases by $\\sim 20\\times 10^{37}$ erg s$^{-1}$  giving $E_{\\rm flare}$ $\\sim 1\\times 10^{41}$ erg which again can be provided by nuclear burning of accumulated matter.  This is completely analogous to nuclear burning in X-ray bursts in which it was realized at an early stage that although the energy release per nucleon is much less than the accretion release, if matter is accumulated and consumed suddenly, the nuclear burning temporarily dominated the energy release (Lamb \\& Lamb 1078).    \\subsection{Does luminosity determine Cyg-like or Sco-like nature ?}  The transient source XTE\\th J1702-462 having an outburst in 2006-7 (Remillard et al. 2006) rose to high luminosity at the start of the outburst and exhibited Cyg-like properties, but as the brightness decreased became more similar to Sco-like sources and eventually displayed Atoll source properties (Homan et al. 2007; Lin et al. 2009). Homan et al. (2010) thus proposed that luminosity, i.e. mass accretion rate determined Cyg-like or Sco-like behaviour, higher $\\dot M$ making the source Cyg-like. As pointed out by Homan et al., this also means that differences in inclination or magnetic field strength cannot determine the nature as previously suggested. The suggestion that luminosity determines Cyg or Sco-like behaviour is not supported by our results as we see no systematic differences of luminosity between the two types as shown in the figures, the Cyg-like sources spanning a factor of three in luminosity and the Sco-like sources almost as much. (It has been previously widely accepted, of course, that Cyg\\th X-2 was less luminous than Sco\\th X-1, for example).  It could be that distance errors lead to the apparent lack of dependence of source type on luminosity. For luminosity to determine type, {\\it all} the Cyg-like sources would have to be more luminous than {\\it all} the Sco-lke sources. In particular, Cygnus\\th X-2 must be more luminous than Sco\\th X-1 for which the distance is well-determined at 2.8$\\pm $0.3 kpc. The distance of Cyg\\th X-2 assumed here to be 9.0 kpc (Table 6) would have to be  $>$ 14.7 kpc for the source to be as luminous as Sco\\th X-1 ($46\\times 10^{37}$ erg s$^{-1}$ at the soft apex).  \\tabcolsep 5.0 mm \\begin{table}[!h] \\begin{center} \\caption{Source distances used in analysis} \\begin{minipage}{84mm} \\begin{tabular}{lr} \\hline \\hline source               &distance\\\\ & (kpc)\\\\ \\noalign{\\smallskip\\hrule\\smallskip} GX\\th 349+2          &9.2\\\\ Scorpius\\th X-1           &2.8\\\\ GX\\th 17+2           &7.5\\\\ Cygnus\\th X-2        &9.0\\\\ GX\\th 340+0          &11.0\\\\ GX\\th 5-1            &9.0\\\\ \\noalign{\\smallskip}\\hline \\end{tabular}\\\\ \\end{minipage} \\end{center} \\end{table}     Distances to Cyg\\th X-2 were published by Cowley et al. (1979) of 5.7 - 8.7 kpc for assumed donor masses of 0.5 - 1.0 M$_{\\sun}$. Orosz \\& Kuulkers derived a donor mass of 0.6$\\pm$0.13 M$_{\\sun}$ and a distance of 7.2$\\pm$1.1 kpc. From peak burst fluxes, Smale (1998) found 11.6$\\pm$0.13 kpc and Galloway (2008) 11$\\pm$2 kpc for an H fraction X = 0.7 and 14$\\pm$3 for X = 0; however, not all bursts exhibited the characteristics of photospheric radius expansion, moreover bursting in the Z-track sources Cyg\\th X-2 and GX\\th 17+2 is not fully \\begin{figure}[!hb]                                                    % \\begin{center} \\includegraphics[width=34mm,height=84mm,angle=270]{fig13a}        % \\includegraphics[width=32mm,height=84mm,angle=270]{fig13b}        % \\caption{ASM lightcurve of the transient XTE\\th J1702-462: \\break upper panel: the complete outburst; lower panel: expanded view showing the limited extent and occurrence of flaring} \\label{} \\end{center} \\end{figure} understood. Jonker \\& Nelemans (2004) give 11.4 - 15.3 kpc for hydrogen rich to helium rich material. Thus the distance is not well-constrained and Cyg\\th X-2 could be {\\it as luminous} as Sco\\th X-1. However, to produce properties very different from Sco\\th X-1 it would have to be appreciably {\\it more luminous} needing a distance more than 20 kpc and this is unlikely.  For GX\\th 340+0, distances have been found of 10.2 kpc (Penninx 1989), 11.8 kpc (van Paradijs \\& White 1995) and 11.0 kpc (Grimm et al. 2002). $L$ is $36.6\\times 10^{37}$ erg s$^{-1}$ at the soft apex for the assumed 11 kpc. This would need to be 12.3 kpc to be as luminous as Sco\\th X-1 which is clearly possible, but $\\sim$17 kpc to be twice as luminous, say, which appears unlikely.  The identification of Sco-like behaviour in the transient was based on the shape of the Z-track. In the notation of Lin et al. (2009) the transient was Cyg-like in section I of the lightcurve and Sco-like in sections II and III. The appearance of hardness-intensity in II was similar to GX\\th 17+2 with a long NB; however, we suggest that this source is not fully Sco-like but transitional between the Cyg and Sco types. In III, the NB becomes small and the FB large and the source does appear Sco-like. However, the intensity increase in flaring was \\hbox{$\\leq$ 40\\%} in the total PCA band while all Sco-like sources have 100\\% increases or more (Figs. 2 and 10), and Cyg\\th X-2 sometimes displays flares of up to 50\\%, so the flaring is not conclusive. In several respects, Homan et al. (2007) also did not find XTE\\th J1702-462 unambiguously Sco-like.   Towards resolving whether the transient source was truly Sco-like, in Fig. 13 the ASM lightcurve of the transient source is shown for comparison with the lightcurves of Sco\\th X-1 and GX\\th 349+2 (Fig. 10). The upper panel shows the complete brightening and decay of the transient; the dotted line shows the part which was thought to be similar to the Sco-like sources; a known modulation at 25 days (Homan et al. 2007) can be seen. In the lower panel, a restricted part of the lightcurve is shown so that the extent of flaring can be seen without this modulation. It is evident that the flaring is not continuous or strong as in Sco\\th X-1 but sparse and weak suggesting that the source was not Sco-like.  Rather than making identification on the appearance of the Z-track, based on the present work we can offer a physical definition of Sco-like nature: that the neutron star blackbody temperature $kT_{\\rm BB}$ is at all times $>$ 2 keV, and that there is frequent strong flaring. Thus GX\\th 17+2 is not fully Sco-like but transitional between Cyg and Sco-like. We suggest that XTE\\th J1702-462 may not have been Sco-like as it decayed but remained Cyg-like with stronger flaring; this is sometimes but not often seen in Cyg\\th X-2. On the basis that the source would move closer to the condition for unstable burning ($\\dot m$ $<$ $\\dot m_{\\rm ST}$) as luminosity decreased, this would be expected. Whether the transient source satisfied the above definition for being Sco-like, or not, could be checked by spectral analysis. If the source did not become Sco-like, then luminosity may not determine whether sources are Cyg-like or Sco-like.      \\subsection{A general Z-track model}  Application of the Extended ADC model having the form discussed in Sect. 1.4 has provided a physical model for the Z-track sources. In summary, in all types of Z-track source, ascending the normal branch is due to an increase of $\\dot M$ for which spectral fitting presents convincing evidence, as a definite increase of both total X-ray luminosity $L_{\\rm Tot}$ and ADC luminosity $L_{\\rm ADC}$. The resultant high neutron star temperature gives a high radiation pressure launching the jets. In Cyg\\th X-2 this predicts an intermittent jet when the source is on that part of the Z. In Sco\\th X-1 the jet may be present even at the soft apex due to the permanently high $kT_{\\rm BB}$.  Flaring in the Cyg-like sources consists of unstable nuclear burning. In the Sco-like sources, $\\dot M$ increases as the source moves away from the soft apex on the NB or FB. Sometimes conditions favour unstable burning so making the FB distinct from the NB. Strong flaring in the Sco-like sources causes the observed spectral differences, i.e. high neutron star temperature and associated radiation pressure. The permanently high $kT_{\\rm BB}$ and high radiation pressure may suggest a permanent bending of the accretion flow out of the plane into the jet. The residual flow to the neutron star may then satisfy the condition for unstable burning ($\\dot m$ $<$ $\\dot m_{\\rm ST}$) so that a period of matter accumulation is followed by flaring, this then repeating. In the Cyg-like sources there is no permanent bending of the accretion flow, so unstable burning only takes place when $\\dot M$ is low. The relevant question then becomes what is the fundamental cause of the strong flaring in the Sco-like sources.  As discussed in Sect. 4.5, the present results do not confirm that Cyg or Sco-like nature is determined by luminosity as suggested by Lin et al. (2009). In this case, some other factor is responsible. Flaring in the Cyg-like sources will only take place when a decrease of $\\dot M$ causes $\\dot m$ to fall below the critical value, and this makes flaring weak and infrequent. The fully Sco-like sources (excepting GX\\th 17+2) are clearly different with longterm diversion of accretion flow into jets implied. It is possible that this could be connected with the long orbital periods of 18.9 and 21.85 hr in Sco\\th X-1 and GX\\th 349+2 as this suggests evolved companions and higher mass accretion rates. This could heat the neutron star causing jet launching and conditions for unstable burning may become more often satisfied.    \\thanks{ We thank the referee for his helpful comments. This work was supported in part by the Polish KBN grants 3946/B/H0/2008/34 and 5843/B/H03/2011/40.}"
    },
    "1208/1208.1666.txt": {
        "abstract": "We study cosmological dynamics and phase space of a scalar field localized on the DGP brane. We consider both the minimally and nonminimally coupled scalar quintessence and phantom fields on the brane. In the nonminimal case, the scalar field couples with induced gravity on the brane. We present a detailed analysis of the critical points, their stability and late-time cosmological viability of the solutions in the phase space of the model.\\\\ {\\bf Key Words}: Braneworld Cosmology, Scalar Fields, Dynamical Systems, Cosmological Viability ",
        "introduction": "In the revolutionary braneworld viewpoint, our universe is a $3$-brane embedded in an extra dimensional bulk. Standard matter and all interactions are confined on the brane; only graviton and possibly non-standard matter are free to probe the full bulk \\cite{1}. Based on the braneworld viewpoint, our universe may contain many more dimensions than those we experience with our senses. The most compelling reasons to believe in extra dimensions are that they permit new connections between physical properties of the observed universe and suggest the possibility for explaining some of its more mysterious features. Extra dimensions can have novel implications for the world we see, and they can explain phenomena that seem to be mysterious when viewed from the perspective of a three-dimensional observer. Even if one is doubtful about string theory due to, for instance, its huge number of landscapes, recent researches have provided perhaps the most compelling argument in the favor of extra dimensions: a universe with extra dimensions might contain clues to physics puzzles that have no convincing solutions without them. This reason alone makes extra dimensional theories worthy of investigation. In this streamline, the braneworld models that are inspired by ideas from string theory provide a rich and interesting phenomenology, where higher-dimensional gravity effects in the early and late universe can be explored, and predictions can be made in comparison with high-precision cosmological data.  Even for the simplest models of Randall-Sundrum (RS) \\cite{2} and Dvali-Gabadadze-Porrati (DGP) \\cite{3}, braneworld cosmology brings new implications on the inflation and structure formation. Also it brings new ideas for dark energy and opens up exciting prospects for subjecting M-theory ideas to the increasingly stringent tests provided by high-precision astronomical observations. At the same time, braneworld models provide a rich playground for probing the geometry and dynamics of the gravitational field and its interaction with matter. In these respects, the DGP braneworld model is a scenario that gravity is altered at immense distances by the excruciatingly slow leakage of gravity off our $3$-brane universe. In this braneworld scenario, the bulk is considered as empty except for a cosmological constant and the matter fields on the brane are considered as responsible for the evolution on the brane \\cite{3}. The self-accelerating DGP branch explains late-time speed-up by itself without recourse to dark energy or other mysterious components \\cite{4}. Even the normal DGP branch has the potential to realize an effective phantom phase via dynamical screening of the brane cosmological constant \\cite{5}.  Scalar fields play a crucial role in modern cosmology, both in models of the early universe and late-time acceleration. Scalar fields provide also a simple dynamical model for matter fields in a braneworld and dark energy models. In the early universe, inflaton as a scalar field provides the required basis of the some well-established inflation models. Also at late time, dark energy models based on dynamical scalar fields have been studied extensively in recent years \\cite{6}. In braneworld models, the existence of a scalar field on the brane provides a variety of possibilities that brings the corresponding theory to explain some novel properties. In fact, a particular form of the bulk or brane matter is a scalar field. In the context of braneworld induced gravity, it is natural to consider a non-minimal coupling of the scalar field and induced Ricci curvature on the brane. The resulting theory can be thought of as a generalization of the Brans\u0096Dicke type scalar-tensor gravity in a braneworld context \\cite{7}. As has been pointed in \\cite{8}, the introduction of the non-minimal coupling (NMC) is not just a matter of taste: the NMC is forced upon us in many situations of physical and cosmological interest. For instance, NMC arises at the quantum level when quantum corrections to the scalar field theory are considered. Even if for the classical, unperturbed theory this NMC vanishes, it is necessary for the renormalizability of the scalar field theory in curved space. In most theories used to describe inflationary scenarios, it turns out that a non-vanishing value of the coupling constant is inevitable. In general relativity, and in all other metric theories of gravity in which the scalar field is not part of the gravitational sector, the coupling constant necessarily assumes the value of $\\frac{1}{6}$ \\cite{8}. Therefore, it is natural to incorporate an explicit NMC between the scalar field and the Ricci scalar in the inflationary paradigm and also in scalar fields models of dark energy. In particular the effect of this NMC in a DGP-inspired braneworld cosmology has been studied by some authors (see \\cite{7} and also \\cite{9}).  There are several studies focusing on braneworld models with brane/bulk scalar fields. Some of these studies concentrate on the bulk scalar fields minimally or nonminimally coupled to the bulk Ricci scalar \\cite{10}. The scalar field minimally or non-minimally coupled to gravity on the brane are studied by some authors \\cite{11,12,13,14}. In \\cite{14}, the authors are studied the self-accelerating solutions in a DGP brane with a scalar field trapped on it within a dynamical system perspective. They have shown that the dynamical screening of the scalar field self-interaction potential occurring within the Minkowski cosmological phase of the DGP model mimics 4D phantom behavior and is an \\emph{attractor} solution for a constant self-interaction potential. But, this is not the case necessarily for an exponential potential. For exponential potential, they have shown that gravitational screening is not even a critical point of the corresponding autonomous system of ordinary differential equations. Along with this pioneer work, we consider a scalar field trapped on the DGP brane and we suppose this scalar field is non-minimally coupled to induced gravity on the brane. We study cosmological dynamics on the normal branch of the scenario within a phase space approach with both quintessence and phantom fields on the brane. We provide a phase space analysis of each model through a detailed study of the fixed points, their stability and cosmological viability of the solutions. We also study the classical stability of the solutions in each case in the $w_{\\varphi}-w_{\\varphi}'$ phase-plane. Our study, in comparison with existing literature in this field, provides a complete framework and contains several aspects of the problem not been considered yet. Since the self-accelerating DGP branch has ghost instabilities, our study here is restricted just to the normal DGP branch of the models. While the normal branch of a pure DGP setup has not the potential to explain the late-time cosmic speed-up and crossing of the phantom divide, we show that with a scalar field on the brane there are several new possibilities in the favor of these observationally supported issues.  We note that our motivation to study this model is as follows: as we have indicated above, observations support (at least mildly) that the equation of state parameter of dark energy has crossed the cosmological constant line ($w=-1$) in recent past (at redshift $z\\sim0.25$). It is impossible to realize this feature with a quintessence or phantom field minimally coupled to gravity in standard 4-dimensional theory \\cite{6}. Although the original DGP model was proposed to realize accelerated expansion of the universe in a braneworld setup, the self-accelerating branch of the DGP cosmological solutions has ghost instability. It is impossible also to cross the phantom divide line without a scalar field in the self-accelerating DGP branch \\cite{11}. The normal DGP solution has no ghost instability, but it cannot explain accelerated expansion and crossing of the phantom divide. It has been shown that localizing a scalar field on the normal DGP setup realizes these features \\cite{11}. It is possible also to incorporate extra degrees of freedom on the braneworld setup to have more successful models (see for instance \\cite{7}. These extra degrees provide new facilities and richer cosmological history on the brane, a part of which is related to the wider parameter space. On the other hand, considering just a cosmological constant on the brane, although explains accelerated expansion through dynamical screening of the brane cosmological constant, it has not the potential to explain crossing of the phantom divide \\cite{5}. In this paper we have shown that a scalar field, minimally or nonminimally coupled to induced gravity on the brane has the potential to fill these gaps. We stress that all of the accelerated phases obtained in this study belong to the normal DGP branch of the model. Our study, based on the phase space analysis, provides the most complete treatment of the issue in the field. We have provided a complete analysis of the problem focusing on all possible details. ",
        "conclusions": "The accelerated expansion of the universe supported by recent observational data could be associated with dark energy, whose theoretical nature and origin are still unknown for cosmologists. Cosmological constant or vacuum energy with an equation of state parameter $\\omega=-1$, is the most popular candidate for dark energy, but unfortunately it suffers from some serious problems such as huge fine-tuning and coincidence problems. Therefore, a number of models containing dynamical dark energy have been proposed as responsible mechanisms for late-time cosmic speed up. Some of these models are quintessence, k-essence, phantom scalar field, chaplygin gas models and so on. Another alternative approach to explain the late-time cosmic speed up is modification of the geometric sector of the Einstein field equations leading to modified gravity theories. In the spirit of modified gravity proposal, the Dvali-Gabadadze-Porrati (DGP) braneworld scenario explains the late-time accelerated expansion in its self-accelerating branch without need to introduce a dark energy component on the brane. However, some important features of dark energy such as possible crossing of the cosmological constant equation of state parameter are missing in the pure DGP model. In addition, the self-accelerating DGP solution suffers from ghost instability which makes the model unfavorable. Incorporation of a scalar field component on the DGP brane and treating the normal branch solutions brings a lots of new physics, some of which are studied in this paper. Previous studies in this field are restricted to either scalar fields on the self-accelerating branch or the simple case of minimally coupled scalar fields. In this paper we considered a scalar field component (both quintessence and phantom scalar fields), non-minimally coupled with induced gravity on the brane. We studied cosmological dynamics of the normal branch solutions on the brane within a dynamical system approach. We translated dynamical equations into an autonomous dynamical system in each case. Then we obtained the critical points of the model in phase space of each model. The issue of stability of these solutions are studied with details. Also possibility of having a stable attractor in de Sitter phase corresponding to current accelerated phase of universe expansion are studied in each case. We have also investigated the possibility to have a transition to the phantom phase of the equation of state parameter in each case. The classical stability of the solutions are treated also in a $w_{\\varphi}-w'_{\\varphi}$ phase-plane analysis in each step. We have shown in each step that there are several phases of accelerated expansion in each case, but only a limited critical points have stable, attractor solutions with de Sitter scale factor describing the current accelerated expansion on the brane. In summary, the main achievements of this study are as follows: \\\\   $\\bullet$ While the pure, normal DGP solution has not the potential to explain late-time cosmic acceleration, with a minimally coupled quintessence field in the normal DGP setup there is a stable de Sitter phase realizing the late-time cosmic speed up. Nevertheless, as the pure DGP case, there is no possibility to cross the cosmological constant line by the effective equation of state parameter of the model. The classical stability domain of the model is restricted to those subspaces of the model parameter space that $w'_{\\varphi}<3w_{\\varphi}(1+w_{\\varphi})$  with $w_{\\varphi}>-1$ where $w'_{\\varphi}\\equiv\\frac{dw_{\\varphi}}{dN}$ and $N\\equiv\\ln a(t)$.\\\\    $\\bullet$ With a minimally coupled phantom field on the brane, it is possible to have an attractor, de Sitter solution realizing the late-time accelerated expansion in the normal DGP setup. Also, the effective equation of state parameter of the model crosses the phantom divide. Note that this crossing is impossible by the effective equation of state parameter of a minimally coupled quintessence field on the brane. Similar to the minimally quintessence field on the brane, the classical stability domain of the model is restricted to those subspaces of the model parameter space that $w'_{\\varphi}<3w_{\\varphi}(1+w_{\\varphi})$ with $w_{\\varphi}<-1$.\\\\   $\\bullet$ With a non-minimally coupled quintessence scalar field on the DGP brane, it is possible to realize a stable, de Sitter late-time accelerating phase in the normal branch of the model. It is possible also to cross the phantom divide by the effective equation of state parameter of the model. In this case, there are two different domains of stability in the $w_{\\varphi}-w'_{\\varphi}$ phase-plane of the model: a subspace with $w'_{\\varphi}>3w_{\\varphi}(1+w_{\\varphi})$ with $w_{\\varphi}<-1$ corresponding to an \\emph{effective phantom phase} and the other subspace with $w'_{\\varphi}<3w_{\\varphi}(1+w_{\\varphi})$ with $w_{\\varphi}>-1$ corresponding to a quintessence phase. This feature shows that it is possible to have an effective phantom picture with a non-minimally quintessence field on the normal DGP setup. \\\\   $\\bullet$ For a non-minimally coupled phantom scalar field on the DGP brane, it is possible to realize a de Sitter late-time accelerating phase in the normal branch of the model. It is possible also to cross the phantom divide line by the effective equation of state parameter of the model in this case. Also, as for the case of non-minimally coupled quintessence field on the brane, there are two different domains of stability in the $w_{\\varphi}-w'_{\\varphi}$ phase-plane of the model: a subspace with $w'_{\\varphi}>3w_{\\varphi}(1+w_{\\varphi})$ with $w_{\\varphi}<-1$ corresponding to the phantom phase and the other subspace with $w'_{\\varphi}<3w_{\\varphi}(1+w_{\\varphi})$ with $w_{\\varphi}>-1$ corresponding to an \\emph{effective quintessence phase} on the brane. Note that it is possible to have an effective quintessence picture with a non-minimally coupled phantom field on the normal DGP setup.  Finally we stress that the observational status of the present DGP-inspired models can be treated in the same line as has been reported in Ref. \\cite{25}."
    },
    "1208/1208.0336_arXiv.txt": {
        "abstract": "Dynamical dark matter (DDM) is an alternative framework for dark-matter physics in which the dark-matter candidate is an ensemble of constituent fields with differing masses, lifetimes, and cosmological abundances.  In this framework, it is the balancing of these quantities against each other across the ensemble as a whole which ensures phenomenological viability. In this paper, we examine the prospects for the direct detection of a DDM ensemble.  In particular, we study the constraints imposed by current limits from direct-detection experiments on the parameter space of DDM models, and we assess the prospects for detecting such an ensemble and distinguishing it from traditional dark-matter candidates on the basis of data from the next generation of direct-detection experiments.  For concreteness, we focus primarily on the case in which elastic scattering via spin-independent interactions dominates the interaction rate between atomic nuclei and the constituent particles of the ensemble.  We also briefly discuss the effects of modifying these assumptions. ",
        "introduction": "}   Dynamical dark matter (DDM)~\\cite{DynamicalDM1,DynamicalDM2} has recently been advanced as an alternative framework for dark-matter physics.  In this framework, the usual assumption of dark-matter stability is replaced by a balancing between lifetimes and cosmological abundances across a vast ensemble of particles which collectively constitute the dark matter.  Within this framework, the dark-matter candidate is the full ensemble itself --- a collective entity which cannot be characterized in terms of a single, well-defined mass, lifetime, or set of interaction cross-sections with visible matter.   As a result, cosmological quantities such as the total relic abundance $\\Omegatot$ of the ensemble, its composition, and its equation of state are time-dependent (\\ie, dynamical) and evolve throughout the history of the universe.  Moreover, for this same reason, DDM ensembles also give rise to a variety of distinctive experimental signatures which serve to distinguish them from traditional dark-matter candidates.  A number of phenomenological and cosmological consequences to which DDM ensembles can give rise were presented in Refs.~\\cite{DynamicalDM2,DynamicalDM3}, along with the bounds such effects imply on the parameter space of an explicit model within the general DDM framework. DDM ensembles can also give rise to characteristic signatures at colliders~\\cite{DDMColliders}, including distinctive imprints on the kinematic distributions of the Standard-Model (SM) particles produced in conjunction with the dark-sector fields.  In this paper, we examine the prospects for the direct detection of DDM ensembles via their interactions with atomic nuclei --- a detection strategy~\\cite{GoodmanWitten} which has come to play an increasingly central role in the phenomenology of most proposed dark-matter candidates (for reviews, see, \\eg, Ref.~\\cite{DirectDetReviews}). Indeed, conclusive evidence of nuclear recoils induced by the scattering of particles in the dark-matter halo would provide the most unambiguous and compelling signal --- and moreover the only non-gravitational evidence --- for particle dark matter to date. Data from the current generation of direct-detection experiments have already placed stringent constraints on many models of the dark sector, and the detection prospects have been investigated for a variety of traditional dark-matter candidates at next generation of such experiments. Studies in the context of particular multi-component models of the dark sector have also been performed~\\cite{MultiComponentBlock,ProfumoTwoComponentDirectDet}.  Here, we shall demonstrate that DDM ensembles can give rise to distinctive features in the recoil-energy spectra observed at direct-detection experiments, and that these features can serve to distinguish DDM ensembles from traditional dark-matter candidates.  These features include resolvable kinks in the recoil-energy spectra, as well as characteristic shapes which are difficult to realize within the context of traditional models --- particularly under standard astrophysical assumptions. As we shall demonstrate, these features should be distinguishable for a broad range of DDM scenarios at the next generation of direct-detection experiments.  Of course, the potential for differentiation within the appropriate limiting regimes of DDM parameter space accords with those obtained in previous studies of two-component models~\\cite{ProfumoTwoComponentDirectDet}. However, as we shall demonstrate, the assumption of a full DDM ensemble as our dark-matter candidate leads to many distinctive features which emerge in significant regions of parameter space and which transcend those which arise for models with only a few dark-sector particles.  This paper is organized as follows. In Sect.~\\ref{sec:ScatteringTradDM}, we review the general aspects of dark-matter direct detection.  We discuss how considerations related to particle physics, nuclear physics, and astrophysics impact both the differential and total rate for the inelastic scattering of dark-matter particles with atomic nuclei, and examine the properties of the recoil-energy spectra associated with traditional dark-matter candidates. In Sect.~\\ref{sec:ScatteringDDM}, by contrast, we investigate how these results are modified when the dark-matter candidate is a DDM ensemble, and we compare the resulting recoil-energy spectra to those obtained in traditional dark-matter models.  In Sect.~\\ref{sec:Limits}, we derive a set of constraints on the parameter space of DDM models from current direct-detection data, and in Sect.~\\ref{sec:Prospects}, we discuss the prospects for obtaining evidence of DDM ensembles at the next generation of direct-detection experiments and for distinguishing such ensembles from traditional dark-matter candidates. Finally, in Sect.~\\ref{sec:Conclusions}, we summarize our results and discuss possible directions for future study. ",
        "conclusions": "}   In this paper, we have investigated the potential for discovering a DDM ensemble and differentiating it from a traditional dark-matter candidate at the next generation of dark-matter direct-detection experiments.  In particular, we have assessed the degree to which these two classes of dark-matter candidates may be distinguished on the basis of differences in recoil-energy spectra. We have demonstrated that DDM ensembles give rise to a number of characteristic features in such spectra, including observable kinks and distinctive ogee profiles.  Moreover, we have demonstrated that under standard assumptions, the identification of such features can serve to distinguish a DDM ensemble from any traditional dark-matter candidate at the $5\\sigma$ significance level at the next generation of direct-detection experiments.  We have found that the prospects for differentiation are particularly auspicious in cases in which the mass splittings between the constituent fields in the DDM ensemble are small and in which the mass of the lightest such field is also relatively small. Note that this is also a regime in which a large fraction of the full DDM ensemble contributes meaningfully to $\\OmegaCDM$.  It is also interesting to compare the prospects for distinguishing DDM ensembles at direct-detection experiments to the prospects for distinguishing them at the LHC.  We have demonstrated here that the former are greatest when $m_0 \\lesssim 20$~GeV, $\\Delta m$ is small, $0.25 \\lesssim\\delta \\lesssim 2$, and the effective couplings between the $\\chi_j$ and SM particles decrease moderately with $m_j$.  By contrast, it was shown in Ref.~\\cite{DDMColliders} that the latter are greatest when $\\Delta m$ is small, $\\delta \\lesssim 1$, and the effective couplings to SM particles increase with $m_j$.  Thus, we see that these two experimental methods of distinguishing DDM ensembles are effective in somewhat different regions of parameter space, and are therefore complementary.  However, we note that there is one region in which evidence for a DDM ensemble may manifest itself both at direct-detection experiments and at the LHC.  This is the region in which $m_0 \\lesssim 20$~GeV, $0.25 \\lesssim \\delta \\lesssim 0.75$, the effective couplings to SM particles are roughly independent of mass, and $\\Delta m$ is also either quite small or else within the range in which an observable kink arises in the recoil-energy spectrum.  The simultaneous observation of both collider and direct-detection signatures in this case would provide highly compelling experimental evidence for a DDM ensemble.  Needless to say, there are numerous additional directions potentially relevant for direct detection which we have not explored in this paper.  For example, we have not considered the prospects for distinguishing DDM ensembles at argon- or carbon-based detectors, or instruments involving target materials other than xenon and germanium.  Likewise, we have not considered the prospects for observing an annual modulation in the signal rate at direct-detection experiments --- a strategy long employed by the DAMA experiment and more recently by CoGeNT.  We have also not considered directional detection.  More generally, we have not considered modifications of the astrophysical assumptions (such as the halo-velocity distributions) or nuclear-form-factor model which define our standard benchmark.  Finally, we have not endeavored to compare or correlate signals from multiple detectors using different target materials.  These directions are all ripe for further study~\\cite{FutureDDMDDPaper}.  In a similar vein, in this paper we have restricted our attention to cases in which elastic processes dominate the scattering rate for all particles in the DDM ensemble.  However, within the context of the DDM framework, inelastic scattering processes~\\cite{InelasticDM,TuckerSmithInelasticDM,MagneticFluffy} of the form $\\chi_j N\\rightarrow \\chi_k N$ where $j\\neq k$ also occur, and can contribute significantly to this rate when $\\Delta m \\lesssim \\mathcal{O}(100\\mathrm{~keV})$. This possibility is particularly interesting for a number of reasons. For example, in DDM scenarios, the final-state particle in such inelastic scattering events can be a heavier particle in the ensemble, as in typical inelastic dark-matter models, but it can also be a {\\it lighter} particle in the ensemble.  In other words, inelastic scattering in the DDM framework involves both ``upscattering'' and ``downscattering'' processes. This latter possibility is a unique feature of DDM scenarios, given that the initial-state particle $\\chi_j$ need not be the lightest particle in the dark sector. Moreover, as we have demonstrated above, the range of $\\Delta m$ relevant for inelastic scattering is also one in which the characteristic features to which DDM ensembles give rise are particularly pronounced.  Some of the consequences of such inelastic processes are readily apparent. For example, let us consider the case in which $|\\delta m_{jk}| \\ll m_j,m_k$, where $\\delta m_{jk} \\equiv m_k - m_j$. Although the matrix element for inelastic scattering is, to leading order, of the same form as for an elastic interaction, the kinematics can be very different.  In the limit $|\\delta m_{jk}| \\ll m_j,m_k$, the maximum recoil energy $E_{jk}^+$ and minimum recoil energy $E_{jk}^-$ possible in inelastic scattering are given by \\begin{equation} E_{jk}^{\\pm} ~=~ \\frac{2 \\mu_{Nj}^2 v^2}{m_N} \\left(1 - \\frac{\\delta m_{jk}}{\\mu_{Nj} v^2} \\pm\\sqrt{1 - 2\\frac{\\delta m_{jk}}{\\mu_{Nj} v^2}} \\right) \\end{equation} where $v$ is the relative velocity of the initial particles.  If $\\delta m_{jk} >0$, then this upscattering process is similar to that typically considered in models of inelastic dark matter~\\cite{InelasticDM,TuckerSmithInelasticDM}, and its basic effect is to narrow the range of recoil energies for which scattering is possible for a fixed dark-matter velocity relative to the Earth.  A general result of this effect is that heavier components of the DDM ensemble are ``brought into range\" of a direct-detection experiment, which can then resolve the recoil-energy endpoint.  For $\\delta m_{jk} <0$, however, the range of possible recoil energies is broadened.  As a result, low-mass members of the DDM ensemble can produce recoils which lie above the recoil-energy threshold $E_R^{\\mathrm{min}}$ of a particular experiment. Moreover, we note that the matrix element for the process $\\chi_j N \\rightarrow \\chi_k N$ determines the matrix element for the process $\\chi_k N \\rightarrow \\chi_j N$ through crossing symmetry.  For a DDM ensemble with a fixed distribution of densities, the scattering rates of different components can thus be related to each other.  All of these possibilities will be discussed further in Ref.~\\cite{DDMDDInelastic}."
    },
    "1208/1208.5513_arXiv.txt": {
        "abstract": "We have conducted three-dimensional self-gravitating radiation hydrodynamical models of gas accretion onto high mass cores (15-33~\\earthmass) over hundreds of orbits. Of these models, one case accretes more than a third of a Jupiter mass of gas, before eventually undergoing a hydrodynamic collapse. This collapse causes the density near the core to increase by more than an order of magnitude, and the outer envelope to evolve into a circumplanetary disc. A small reduction in the mass within the Hill radius (\\rhill) accompanies this collapse as a shock propagates outwards. This collapse leads to a new hydrostatic equilibrium for the protoplanetary envelope, at which point 97 per cent of the mass contained within the Hill radius is within the inner 0.03~\\rhill \\ which had previously contained less than 40 per cent. Following this collapse the protoplanet resumes accretion at its prior rate. The net flow of mass towards this dense protoplanet is predominantly from high latitudes, whilst at the outer edge of the circumplanetary disc there is net outflow of gas along the midplane. We also find a turnover of gas deep within the bound envelope that may be caused by the establishment of convection cells. ",
        "introduction": "The ideas discussed in this paper begin with \\cite{PerCam1974}, who stated that ``when the mass of the core becomes sufficiently great, the surrounding gaseous envelope will become hydrodynamically unstable against collapse onto the planetary core''. This process is controlled by the battle between gravity that acts to contract an envelope onto the core and the gas pressure which acts to support it.  \\cite{Miz1980} performed stability calculations to determine the combinations of core masses and opacities that would make a protoplanetary envelope unstable to collapse. Work of his contemporaries considering the structure of giant planets in the solar system suggested that each such planet (Jupiter, Saturn, Uranus, \\& Neptune) possessed a solid core with a mass of order 10~\\earthmass \\ \\citep{Sla1977,HubMac1980}. Using these values as a target, and assuming a fixed accretion rate, \\citeauthor{Miz1980} concluded that a grain opacity of $\\kappa \\approx 1~{\\rm cm^2 \\: g^{-1}}$ was required in the envelope material during formation to trigger a collapse when the core mass was $\\approx 10$~\\earthmass. Lower opacities were found to lead to envelope collapse at lower core masses. Following the envelope collapse \\citeauthor{Miz1980} states that continued accretion is likely required for the protoplanet's to attain their final masses.  It was suggested later by \\cite{BodPol1986}, who performed evolution calculations that included evolution beyond the attainment of the critical core mass, that rather than a dynamical collapse, the envelope may quasi-statically contract onto the solid core. They suggested that if a sufficient mass of molecular hydrogen in the envelope was apt to undergo dissociation, then this would remove enough energy from the contraction to bring about a dynamical collapse. However, their models indicate that such dissociative regions possess an insufficient fraction of the envelope mass for this to occur.  In the early 1990s \\citeauthor{Wuc1990} wrote a series of papers exploring the evolution of a protoplanetary envelope through the hydrostatic phases approaching the critical core mass, and in what he found to be a subsequent hydrodynamic phase \\citep{Wuc1991}. Solving the equations of radiation hydrodynamics in one-dimension, \\citeauthor{Wuc1991} indicated that following a period of quasi-static contraction, during which the envelope heats up, the transport of this heat out through the envelope by convection and radiation perturbs the hydrostatic equilibrium. In particular \\citeauthor{Wuc1991} cites the $\\kappa-\\rm{mechanism}$ as the means of exciting the dynamical waves that destabilise the envelope. The result of this hydrodynamic evolution was the ejection of a large fraction of the envelope mass, rather than an inwards collapse as had previously been supposed by \\cite{PerCam1974}.  \\cite{PolHubBodLis1996} continued performing models assuming a quasi-hydrostatic contraction, and suggested that further work was required to consider the hydrodynamic evolution of young protoplanets to establish the proper evolution scenario. A step in this direction was taken by \\cite{TajNak1997} who performed a stability analysis of a growing protoplanetary envelope using a distinct numerical code to that of previous authors. Their aim was to determine whether quasi-static contraction, or an envelope instability akin to that suggested by \\citeauthor{Wuc1991}, was the more likely evolutionary course for a growing protoplanet. They perturbed the envelope at intervals during its evolution to see if such action might push a marginally stable system towards instability. They concluded that quasi-static contraction was viable for a protoplanet growing all the way to a Jupiter mass.  There has since been a substantial amount of work considering the gas accretion rates that cores might achieve in circumstellar discs with a variety of properties. \\cite{IkoNakEmo2000} performed quasi-static evolutionary models to determine the dependence of gas accretion upon core mass, grain opacity, and the core's accretion history, finding that these factors were strongly, moderately, and weakly significant respectively. \\cite{BryCheLinNel1999} and \\cite{LubSeiArt1999} performed locally-isothermal hydrodynamics simulations of discs containing planetary cores, the latter finding that the accretion rate drops off as the protoplanet mass becomes very large; a result of the broadening disc gap that it forms. Further hydrodynamic models with more realistic thermodynamics were performed by \\cite{DAnHenKle2003}, \\cite{KlaKle2006}, \\cite{PaaMel2008}, and \\cite{AylBat2009}, finding similar turn overs in the accretion rate with increasing mass, and illustrating the impact of grain opacity on accretion. These models have generally had limited resolution in the vicinity of the protoplanet, and though \\cite{AylBat2009} achieved high resolution, the evolutionary period was extremely short due to the computational demands of the calculations.  In this paper we report results from three-dimensional self-gravitating radiation hydrodynamics calculations that resolve the protoplanet's envelope, whilst modelling its hydrodynamic evolution and growth within a section of a circumstellar disc. Using high mass discs (though still stable; Toomre Q $>> 1$), and assuming low opacities, we achieve accretion rates that allow significant envelope growth in only a few hundred orbits, allowing us to examine their development. Our computational method is described in Section~\\ref{sec:setup}, followed by our results in Section~\\ref{sec:results}, and a discussion of their relationship with previous results in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}   \\subsection{Hydrodynamic collapse}  From the earliest suggestion of \\cite{PerCam1974} it has been thought that a giant planet might form through the hydrodynamic collapse of a gaseous envelope onto a solid core which caused it to assemble. This was followed by numerous models that effectively sought for hydrostatic solutions to various combinations of properties to establish when such a collapse might occur \\citep{MizNakHay1978,Miz1980,Sas1989}. The first models that attempted to model giant planet growth from the initial core formation, through to the envelope growth were performed by \\cite{BodPol1986}. These models revealed that a protoplanetary envelope would gradually contract as the planet grew, leading to a quasi-static contraction beyond previously calculated values for the critical mass, as long as the envelope did not effectively detach from the protoplanetary disc (that is there was a sufficiently rapid supply of material from the latter to the former). Under these conditions, their was no evidence to suggest that the hydrostatic balance should reach some limit beyond which a collapse was inevitable, and later semi-analytic models originating from these earlier works, such as \\cite{PolHubBodLis1996}, \\cite{HubBodLis2005}, and \\cite{LisHubDAnBod2009}, suggest no need for a dynamic collapse. Our Model J follows a pattern of stable growth for the vast majority of its history, though not evidently undergoing any significant contraction, and resumes this pattern of growth subsequent to its envelope collapse.   \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{figure13.pdf} \\caption{Velocity vectors in the plane of the disc, illustrating the gas flow about the Hill radius (marked with a dashed line) that leads to the alternating pattern of in and out flow seen in the first panels of the mass flux plots shown in Figs.~\\ref{fig:preflow} \\& \\ref{fig:postflow}. This vector field is plotted for a time preceding the envelope collapse, but a very similar field exists after the shock associated with the collapse has passed out of the region.} \\label{fig:vfield} \\end{figure}    The models presented in this article are performed using a three-dimensional hydrodynamics code that include self-gravity, and radiative transfer, but which omits the core formation phase, and the deposition of energy due to planetesimal accretion that are included in the semi-analytic works discussed. However, at the time of interest around the envelope collapse, the energy release is utterly dominated by the contraction of the gaseous envelope, such that solids accretion energy may be regarded as negligible. The metallicity of the envelope might be significantly modified by the ablation and evaporation of grains that have been accreted over the planets history, and we make no attempt to account for this. The opacity in Model J is reduced by a fixed factor of $10^{3}$, { at the lowest end of the range suggested by \\cite{MovBodPodLis2010}, who found such opacities due to grain settling and coagulation in regions of the envelope.}  It is difficult to disentangle the causes and effects of the very rapid collapse we find, for example the surge in temperature leads to a higher fraction of dissociated hydrogen. However, as stated in Section~\\ref{sec:collapse}, it does not appear that the dissociation of molecular hydrogen acts to trigger the collapse, as the fraction remains steady in the preceding period. There is also no evidence of the Kappa-mechanism acting within the protoplanetary atmosphere in our model, as was found by \\cite{Wuc1991} to cause a dynamic collapse. \\citeauthor{Wuc1991} also found that this collapse led to a significant ejection of material from the protoplanet, whilst we find only a small drop of $\\approx 1.3$ per cent in the mass within the Hill radius, and a rapid increase at and below 0.1~\\rhill \\ as the protoplanet structure shifts to its new state.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{figure14.pdf} \\caption{Each panel shows the ratio of the pressure force to the gravitational force in the envelope and beyond, with the first and last panels corresponding in time to the pre and post-collapse cumulative mass distributions shown in Fig.~\\ref{fig:cmass}. A ratio of one indicates that the material is in hydrostatic equilibrium, which before the envelope collapse applies to a region out to 0.0067~\\rorbit \\ (0.12~\\rhill, top panel, marked with vertical dotted line), but which post-collapse reaches out to just 0.0009~\\rorbit \\ (0.016~\\rhill). By the final panel the shock wave has cleared the inner 0.01~\\rorbit, leaving the environment internal to this in a new hydrostatic equilibrium, whilst its mass continues to increase, as shown in Fig.~\\ref{fig:racc}.} \\label{fig:hydrocollapse} \\end{figure}   It appears that our Model J protoplanet reaches the hydrostatic limit for its formative structure, and that the internal pressure gradient can no longer accommodate the addition of mass by a small adjustment. This is illustrated in Fig.~\\ref{fig:hydrocollapse} which shows the ratio of the pressure force to the gravitational force against radius at a number of stages of the envelope collapse. From an initial state, in which the protoplanetary atmosphere is in hydrostatic equilibrium out to a radius of 0.0067~\\rorbit \\ (0.12~\\rhill), the atmosphere rapidly begins to restructure, pushing the hydrostatic region down to a radius of $\\approx 0.0013$~\\rorbit \\ (0.025~\\rhill). This initial stage leaves the form of the graph otherwise relatively unchanged, but as the central concentration of mass continues, a shock begins to form as the pressure near the core surges. The maximum pressure within 0.001~\\rorbit \\ has increased by an order of magnitude between the first panel and the third, leading to a somewhat steeper gradient over this region. However, at this point in time the gradient between $0.001 - 0.002$~\\rorbit \\ steepens much more rapidly, forming a shock front, and this front marks the new radial limit of hydrostatic equilibrium as can be seen in the third panel. The subsequent panel shows the shock propagating outwards, whilst the hydrostatic core shrinks a little. By the final panel the inner envelope has resettled and the structure has stabilised, whilst the shock's propagation continues outwards, eventually moving beyond the limits of the modelled region.  It is possible that this envelope collapse only occurs due to the high accretion rates achieved due to our selected disc conditions. The low opacity assumed promotes very rapid planet growth, and it may be this rapidity that prevents the envelope from adjusting its structure more gradually to accommodate the increasing mass. As such, it may be that the hydrodynamical collapse of a planetary atmosphere can only occur if that planet if accreting very rapidly. Further models will be required to determine whether or not this is the case. We note however that the accretion rate measured in Model J, both just before and after the collapse, is still not as rapid as would be found using a locally-isothermal equation of state, despite the large reduction in opacity. \\cite{LisHubDAnBod2009} present accretion rates in their fig.~3, where these rates were obtained by \\cite{DAnKleHen2003} using three-dimensional hydrodynamical models with a locally-isothermal equation of state. Applying our disc conditions to their results yields an accretion rate of $8 \\times 10^{-4} ~{\\rm M_{Jupiter}~year^{-1}}$, which is twice the rate measured in our radiation hydrodynamics { models prior to collapse}.  { At this juncture we note that the results given in \\cite{LisHubDAnBod2009} show a viscosity dependence, where higher viscosities lead to more rapid accretion. In our calculations viscosity is not a constant, but is proportional to the spatial resolution of the SPH method. Thus, the viscosity is lower in regions of higher density, and these differences mean the above comparison is only approximate. Further, we note that in the absence of a protoplanet, the unperturbed circumstellar disc in our models has a viscosity of $\\alpha \\approx 4 \\times 10^{-3}$, consistent with the fixed viscosity global models of \\citet{BatLubOgiMil2003} that are used to inject gas at the boundaries of the disc section. It is these boundaries that determine the rate at which gas is supplied to the disc section in these local models. As such, once the disc is perturbed, the spatially varying viscosity of the SPH calculations leads to to an inconsistency with the global models, and so the boundaries. A further caveat arising from the boundary implementation is that the injected material comprises gas on both circulating and librating orbits \\citep{LubSeiArt1999}, orbits that are modified as the gas passes through the modelled disc section. However, these modifications are lost when the gas leaves the section, and new gas is injected without these modifications, leading to a further inconsistency.}   \\subsection{Atmospheric turn over} \\label{sec:convection}  As briefly mentioned in Section~\\ref{sec:flow} in reference to the third panel of Fig.~\\ref{fig:postflow}, there appears to be significant motion of the bound gas besides rotation in the plane of the disc. An apparent rolling motion is indicated by the flow through the surface at 0.03~\\rhill, the gas unable to escape, but flowing out and in within the bound atmosphere. These flows appear to be convection cells in the deepest parts of the planetary envelope where the medium is extremely optically thick, and thus unable to cool readily by radiation. { Fig.~\\ref{fig:convection} shows a region out to $0.07$~\\rhill \\ ($3.7 \\times 10^{-3}$~\\rorbit), at which radius the temperature is 675K, increasing to a maximum of 7400K near the core, demonstrating a significant thermal gradient against radius.} \\cite{BodPol1986} found strata of convection within protoplanetary envelopes in their one-dimensional models, with a dense inner convection region that contained some 95 per cent of the envelope mass, and depending on the age and core size, multiple higher level convection zones were also found to form. However, it is likely that the convection in fully 3-D models that include rotation, is very different.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{figure15.pdf} \\caption{A slice through Model J after the envelope has collapsed, showing the mean velocity vector field over the course of 4 orbits which correspond to those used to produce the mass flux plots in Fig.~\\ref{fig:postflow}. The outer circle corresponds to the spherical surface of Fig.~\\ref{fig:postflow} at 0.03~\\rhill, whilst the inner filled circle illustrates the radius of the planetary core. This slice in the z-${\\rm r_{cyl}}$ plane rotates about the core at the same average angular velocity as the gas, which exhibits solid-body rotation within 0.04~\\rhill, so that the vector field is seen in the rotating frame of the gas. The bold sections of the outer circle cover the latitudes of $45-70$ degrees at which out flow was seen in the final panel of Fig.~\\ref{fig:postflow}.} \\label{fig:convection} \\end{figure}  Fig.~\\ref{fig:convection} illustrates the vector field in a slice through the protoplanet, averaged over 4 orbits that correspond to those used to produce the mass flux rendering of Fig.~\\ref{fig:postflow}. Within $\\sim 0.04$~\\rhill \\ the protoplanet rotates as a solid body, with a rate of $4.65 \\times 10^{-7}~{\\rm rad \\: s^{-1}}$ (giving a day equivalent to $\\approx 160$ Earth days). The plane to which the velocity vectors have been mapped was rotated at this rate, such that the inner region of the atmosphere remained consistently aligned with the plane over the time of averaging. There are four distinct regions of turn over in which gas flows outwards at mid-latitudes before falling inwards again along the disc midplane. These looping flows are subsonic, with mean velocities of less than $3~{\\rm m \\: s^{-1}}$ and a maximum of less than $8~{\\rm m \\: s^{-1}}$ in a region where the sound speed is greater than $6~{\\rm km \\: s^{-1}}$. A time average to a fixed plane, or a spatial average rotating about the z-axis at a single moment in time both reveal a similar structure, suggesting that the general form of these structures is persistent. The bold sections of the 0.03~\\rhill \\ circle are marked between $45 - 70$~degrees, the latitudes where outflow was seen in the fourth panel of Fig.~\\ref{fig:postflow}. As might be expected, the velocity vectors across these bold segments indicate outflow, illustrating that these looping flows are responsible for the features in the rendered plot. The 0.03~\\rhill \\ region marked by the outer circle contains $\\approx 97$ per cent of the total protoplanetary mass, that is the mass measured out to the Hill radius; prior to collapse the same region had contained just under 40 per cent of the mass within the Hill radius. { In this dense environment, possessing a significant radial temperature gradient, the 4 distinct cells revealed in the velocity field might well be indications of convection.}  \\begin{figure*} \\centering \\subfigure % { \\includegraphics[width=7cm]{figure16a.pdf} } \\subfigure % { \\includegraphics[width=7cm]{figure16b.pdf} \\vspace{-6pt} } \\subfigure % { \\includegraphics[width=7cm]{figure16c.pdf} } \\subfigure % { \\includegraphics[width=7cm]{figure16d.pdf} }  \\caption{Longitudinally (azimuthally) integrated mass flux for Model J at the same four radii considered in Figs.~\\ref{fig:preflow} \\& \\ref{fig:postflow}. The solid lines represent the pre-collapse case (Fig.~\\ref{fig:preflow}) and the dashed lines the post-collapse case (Fig.~\\ref{fig:postflow}). In the former, it is notable that the only instance in which an outward flux can be seen is for the largest radii, where this is easily understood by consideration of the large scale flows (spiral shocks, horseshoe regions, and vertical accretion) surrounding the planet, and persists in a similar form post-collapse. Post-collapse there is more structure visible in the flow at small radii. At 0.3~\\rhill \\ there is a midplane outflow that corresponds to that visible the corresponding panel of Fig.~\\ref{fig:postflow}, whilst at 0.1~\\rhill \\ all material is flowing inwards, though at the midplane the flux is small relative to other latitudes. Finally, at 0.03~\\rhill \\ there is evidence of significant turnover in the bound envelope post-collapse, possibly indicative of convection.} \\label{fig:thetaboth} \\end{figure*}     \\subsection{Gas accretion}  \\cite{TanOhtMac2012} have recently examined gas flow onto and within circumplanetary discs. They found, in agreement with the work of \\cite{MacKokInuMat2008}, that gas flowed onto a circumplanetary disc predominantly in the vertical direction. Moreover, their works suggests that material is flowing out along the midplane of a circumplanetary disc. Once the envelope collapses in Model J of this work, and a circumplanetary disc forms, we also find that material is flowing outwards along the midplane beyond a radius of 0.17~\\rhill, which is at around half the outer radius of the circumplanetary disc \\citep{QuiTri1998, AylBat2009a, MarLub2011}. Fig.~\\ref{fig:thetaboth}, which is equivalent to Fig.~6 in \\cite{TanOhtMac2012}, illustrates the longitudinally (equivalently, azimuthally) integrated mass flux for the pre (solid line) and post (dashed line) collapse flows shown in Figs.~\\ref{fig:preflow} \\& \\ref{fig:postflow}. Before the envelope collapses, at 0.3~\\rhill \\ the outflow seen in two places at the midplane is counterbalanced by the associated midplane inflows and the inflow at higher latitudes, yielding a net inflow, as shown by the solid line in the second panel of Fig.~\\ref{fig:thetaboth}. However, post collapse the consistent midplane outflow seen for this radius at all longitudes results in a peak of outflowing material in a region $\\pm 30$~degrees latitude.  About the planet, at one Hill radius the flow is dominated by the streams of material passing in and out of the region due to the form of their circumstellar orbits (see Fig.~\\ref{fig:vfield}). As such, it is unsurprising that the fluxes we obtain, and those of \\cite{TanOhtMac2012} scaled to our model, are similar at this radius; note we make the following comparisons using our post-collapse case which better resembles \\citeauthor{TanOhtMac2012}'s models. At 1~\\rhill \\ they obtain a peak outflow along the midplane of $1.8 \\times 10^{-7}~{\\rm g \\: cm^{-2} \\: s^{-1}}$, whilst we obtain a value of $2.5 \\times 10^{-7}~{\\rm g \\: cm^{-2} \\: s^{-1}}$. The form is also similar, with wings of inflow at higher latitudes of similar magnitude; at the highest latitudes, our normalisation by area leads to a non-zero inflow. However, at smaller radii the mass fluxes we find are considerably larger than at the Hill radius, which differs substantially from \\citeauthor{TanOhtMac2012}'s results, where the peak flux at all radii differ by less than an a factor of 4. At 0.3~\\rhill \\ the peak outward flux has grown to $2.2 \\times 10^{-6}~{\\rm g \\: cm^{-2} \\: s^{-1}}$, and at 0.1~\\rhill \\ the outflow has ceased, but the inflow flux at high latitudes has increased by another factor of around 5. The final panel of Fig.~\\ref{fig:thetaboth} reveals the mass flow resulting from the formation of convection cells in the deep atmosphere post collapse."
    },
    "1208/1208.5039_arXiv.txt": {
        "abstract": "Recent observations show evidence that high-z ($z\\sim 2 - 3$) early-type galaxies (ETGs) are more compact than those with comparable mass at $z\\sim 0$. Such a size evolution is most likely explained by the `Dry Merger Sceanario'. However, previous studies based on this scenario are not able to consistantly explain both the properties of the high-z compact massive ETGs and the local ETGs. We investigate the effect of multiple sequential dry minor mergers on the size evolution of the compact massive ETGs. From an analysis of the Millennium Simulation Database, we show that such minor (stellar mass ratio $M_{2}/M_{1} < 1/4$) mergers are extremely common during hierarchical structure formation. We perform N-body simulations of sequential minor mergers with parabolic and head-on orbits, including a dark matter component and a stellar component. Typical mass ratios of the minor mergers are $1/20 < M_{2}/M_{1} \\lid 1/10$. We show that sequential minor mergers of compact satellite galaxies are the most efficient at promoting size growth and decreasing the velocity dispersion of the compact massive ETGs in our simulations. The change of stellar size and density of the merger remnants is consistent with recent observations. Furthermore, we construct the merger histories of candidates for the high-z compact massive ETGs using the Millennium Simulation Database, and estimate the size growth of the galaxies by the dry minor merger scenario. We can reproduce the mean size growth factor between $z=2$ and $z=0$, assuming the most efficient size growth obtained during sequential minor mergers in our simulations. However, we note that our numerical result is only valid for merger histories with typical mass ratios between 1/20 and 1/10 with parabolic and head-on orbits, and that our most efficient size growth efficiency is likely to an upper limit. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0793_arXiv.txt": {
        "abstract": "The observed amount of lithium for low metallicity population II stars (known as the Spite plateau) is a factor of $\\sim 3-5$ lower than the predictions of the standard cosmology. Since the observations are limited to the local Universe (halo stars, globular clusters and satellites of the Milky Way) it is possible that certain physical processes may have led to the spatial separation of lithium and local reduction of [Li/H]. We study the question of lithium diffusion after the cosmological recombination in sub-Jeans dark matter  haloes, taking into account that more than 95\\% of lithium remains in the singly-ionized state at all times. Large scattering cross sections on the rest of the ionized gas leads to strong coupling of lithium to protons and its initial direction of diffusion coincides with that of H$^+$. In the rest frame of the neutral gas this leads to the diffusion of H$^+$ and Li$^+$ out of overdensities with the trend of reducing [Li/H] in the minima of gravitational wells relative to the primordial value. We quantify this process and argue  that, with certain qualifications, it may have played a significant role in creating local lithium deficiency within the primordial dark matter haloes, comparable to those observed along the Spite plateau. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.5380_arXiv.txt": {
        "abstract": "{} {This work aims at investigating the molecular gas in the surroundings of the ultra-compact HII region G045.47+0.05 looking for evidence of molecular outflows.}{We carried out observations towards a region of 2\\m$\\times$2\\m~centered at RA=19$^{\\rm h}$14$^{\\rm m}$25.6$^{\\rm s}$, dec.=+11\\d09\\m27.6\\s (J2000) using the Atacama Submillimeter Telescope Experiment (ASTE; Chile) in the $^{12}$CO J=3--2, $^{13}$CO J=3--2, HCO$^+$ J=4--3 and CS J=7--6 lines with an angular resolution of 22\\s. We complement these observations with public infrared data.} {We characterize the physical parameters of the molecular clump where G045.47+0.0 is embedded. The detection of the CS J=7--6 line emission in the region reveals that the ultra-compact HII region G045.47+0.0 has not completely disrupted the dense gas where it was born. The HCO$^+$ abundance observed towards G045.47+0.0 suggests the presence of molecular outflow activity in the region. From the analysis of the $^{12}$CO J=3--2 transition we report the presence of bipolar molecular outflows with a total mass of about 300~\\msol. We derive a dynamical time (flow's age) of about $10^5$ yr for the outflow gas, in agreement with the presence of an ultra-compact HII region. We identify the source 2MASS 19142564+1109283 as the massive protostar candidate to drive the molecular outflows. Based on the analysis of its spectral energy distribution we infer that it is an early B-type star of about 15~\\msol. The results of this work support the scenario where the formation of massive stars, at least up to early B-type stars, is similar to that of low mass stars.}  {}   \\titlerunning{Molecular outflows in G045.47+0.05} \\authorrunning{M. E. Ortega} ",
        "introduction": "The formation of high-mass stars remains one of the most significant unsolved problems in astrophysics.  Despite the importance that massive stars have in the structure and dynamic of the Galaxy, the physical processes involved in their formation are less understood than those of their low-mass counterpart.  Observationally, the main problems arise from the fact that they are heavily obscured by dust, are rare, and evolve very fast, making their detection very difficult. The study of massive star formation also poses major theoretical challenges because they begin burning their nuclear fuel and radiating prodigious amounts of energy while still accreting. If the formation of massive stars is similar to that of low mass stars (i.e. via accretion from the surrounding envelope), a mass accretion rate $\\dot{M}$ of at least several orders of magnitude above the values appropriate for the low mass star formation is required \\citep{tan02}.  At present, two theoretical scenarios are proposed to explain the formation of these stars: a monolithic collapse of turbulent gas on the scale of massive dense cores \\citep{tan02}, which is a scaled-up version of the low-mass star formation picture, and a competitive one where accretion occurs inside the gravitational potential of a cluster-forming massive dense core \\citep{bon06}. This last model predicts that stars located near the center of the full gravitational potential accrete at much higher rates than do isolated stars.  The detection of molecular outflows associated with both, high and low-mass young stellar objects, supports a scaled-up formation picture \\citep{beu02}.  In this context, the study of massive molecular outflows and the associated driving source might contribute to discern which is the scenario that prevails in a given star forming region.   It is well known that the generation of molecular outflows during the formation of a high-mass star is a phenomenon that can take place even when the UCHII region stage has been reached (\\citealt{hunter1997}; \\citealt{qin08}). In this work we report the study of the UCHII region G045.47+0.05 (hereafter G45.47), through the analysis of its associated molecular gas and searching for evidence of molecular outflows and its associated driving source. G45.47 was first detected by \\citet{wood89} in radio continuum at 6~cm. This object is adjacent to the extensively studied UCHII region G45.45+0.06 (hereafter G45.45), which is part of a complex of five radio compact HII regions (\\citealt{mat77}; \\citealt{giv05a}; \\citealt{giv05b}). Such complex and G45.47 are embedded in the molecular cloud GRSMC G045.49+00.04 (at V$_{LSR} \\sim$ 58~\\k; \\citealt{rat09}) and are located on the northern border of the more extended HII region named G45L in \\citet{paron09}. Based on HI absorption profiles, \\citet{kuc94} derived a kinematics distance of 8.3~kpc for the UCHII region G45.45. As G45.47 is part of the same complex, we adopt for this object the same distance.   \\citet{cas95} detected a class II CH$_3$OH maser emission at 6.6~GHz at V$_{LSR}\\sim56$~\\k~ towards G45.47. Given that the 6.6~GHz methanol maser is radiatively pumped by IR emission from the warm dust associated exclusively with massive young stellar object (MYSOs)\\footnote{We define MYSOs to be young stellar objects (YSOs) that will eventually become main-sequence O or early B type stars (M$_*$ $\\geq$ 8~\\msol).} their detection is useful to study the kinematic of the gas and, in particular, to establish the systemic velocity \\citep{cyga09}.  Based on high resolution molecular line observations towards G45.47, \\citet{wil96} identified five HCO$^+$ (1--0) clumps and suggested that G45.47 is in the early stages of the formation of an OB cluster. From an ammonia absorption study, \\citet{hofner99} suggested the presence of a remnant molecular core infalling onto the UCHII region. Later, \\citet{cyga08} identified an ``Extended Green Object'' (EGO) at the position of G45.47, the EGO G45.47+0.05. Their identification as EGOs comes from the common coding of the 4.5~$\\mu$m band as green in the three-color composite Infrared Array Camera (IRAC; \\citealt{fazio04}) images from the {\\it Spitzer} Telescope. Extended 4.5~$\\mu$m emission is thought to evidence the presence of shocked molecular gas in protostellar outflows. The association of EGOs with IRDCs and 6.7~GHz CH$_3$OH maser suggests that EGOs trace the formation of massive protostars. According to \\citet{cyga08}, an EGOs is a MYSO with ongoing outflow activity and actively accreting.  In summary, G45.47 is a rich and complex region to study the formation of a new generation of massive stars.  In this paper we present a new study of the dense ambient medium where the UCHII region is evolving. We investigate the molecular gas through several molecular lines observed with the Atacama Submillimeter Telescope Experiment (ASTE; Chile) and characterize the central source based on infrared public data. ",
        "conclusions": "Figure \\ref{intro}-(left) shows a composite {\\it Spitzer}-IRAC three-color image (3.6~$\\mu$m = blue, 4.5~$\\mu$m = green, and 8~$\\mu$m = red). The extended HII region G45L is centered at RA = 19$^{\\rm h}$14$^{\\rm m}$17$^{\\rm s}$, dec.=+11\\d07\\m48\\s (J2000) and is delimited by two arclike structures observed at 8~$\\mu$m.  The UCHII region G45.47 (a.k.a. EGO G045.47+0.05) is the green source located at RA=19$^{\\rm h}$14$^{\\rm m}$25.6$^{\\rm s}$, dec.=+11\\d 09\\m 27.6\\s (J2000). The white box indicates the region mapped with ASTE. A zoom up of this region (Fig. \\ref{intro}-(right)) shows that the UCHII region is embedded in the {\\it Spitzer} dark cloud SDC G45.467+0.048 \\citep{peretto09}. Infrared dark clouds (IRDCs) are dense molecular clouds which appear as extinction features against the bright mid-infrared Galactic background and have been suggested as the cold precursors to high-mass stars \\citep{rath06}.   \\subsection{The molecular gas} \\label{molecular}   In this section, we present the molecular results starting the description with the transitions that trace the most inner part of the clump related to G45.47, and moving to those that map the external layers which may give information about the dynamic effects occurring in the gas.  Figure \\ref{CSspectra}-a shows the CS J=7--6 spectra obtained towards the 2\\m~$\\times$ 2\\m~region (white box in Fig. \\ref{intro}-left). The mapped area includes G45.47 and part of the nearby UCHII region G45.45, located at the (0, 0) and ($-$60, $-$20) offset, respectively. As expected, the most prominent emission arises from these two position, since the detection of the CS J=7--6 transition reveals the presence of warm and dense gas. Figure \\ref{CSspectra}-b shows an enlargement of the spectrum observed at (0, 0). The emission related to G45.47 shows a triple peak structure with velocity components centered at about 56 (the systemic velocity of the gas), 62, and 65~\\k (see Table \\ref{gaussianfit}) with a pronounced dip at about 59 and a less conspicuous one at $\\sim $64~\\k. The velocity component centered at 56~\\k is weak, with a peak about 3$\\sigma$ of the rms noise level. Since the probability of superposition along the line of sight of more than one component in the CS J=7--6 line is very low, we suggest that the dips reveal self-absorption effects in the gas, consistent with an optically thick transition. Self-absorption features demonstrate the existence of a density gradient in the clump \\citep{hira07} and is a common feature of optically thick lines in the direction of embedded young stars. The depression discloses the presence of relatively cold, foreground gas that absorbs photons from warmer material behind it. The CS J=7--6 profile is asymmetric with respect to the dip near v=59~\\k as the redshifted component is clearly brighter than the blueshifted one, suggesting that the molecular gas is expanding. This is because in an expanding cloud a line emission is composed by red and blueshifted photons, the redshifted photons will encounter fewer absorbing material (which is expanding outward) than would blueshifted photons and hence have greater probabilities of escaping (e.g. \\citealt{leu78}; \\citealt{zhou92}; \\citealt{leh97}).  Figure \\ref{CS-maps} shows the velocity channel maps of the CS J=7--6 emission averaged every 1~\\k. It can be noticed the presence of molecular gas associated with G45.47 (red cross) in the velocity range going from $\\sim$ 55 to $\\sim$65~\\k. The other most conspicuous molecular condensation, partially observed, corresponds to the gas related to G45.45 and has a central velocity of $\\sim$ 59~\\k.  \\begin{figure*} \\centering \\includegraphics[width=14cm]{CSspectra.eps} \\caption{a)CS J=7--6 spectra obtained towards the 2\\m~$\\times$~2\\m~region (white box in Fig. \\ref{intro}-left) mapped with ASTE. b) Hanning smoothed profile of the CS J=7--6 line towards G45.47 at the position (0, 0).  The dashed line indicates the systemic velocity of the gas. The spectrum was deconstructed using three Gaussians, which are shown in red.} \\label{CSspectra} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=12cm]{CS-maps.ps} \\caption{Velocity channel maps of the CS J=7-6 emission averaged every 1~\\k.  Greyscale goes from 0.15 (about 3$\\sigma$ rms noise of an averaged channel map) up to 1.8 K. Contours levels are at 0.15, 0.3, 0.45 and 0.6 K. The red cross represents the position of G45.47.} \\label{CS-maps} \\end{figure*}  Figure \\ref{HCO_spectra}-a shows the HCO$^+$ J=4--3 spectra obtained towards the same region. As in the case of the CS J=7--6 spectrum it can be seen that the HCO$^+$ emission towards (0, 0) can be deconstructed in three velocity components, a very weak one centered at about 56~\\k~(intensity of about 2$\\sigma$), plus two brighter ones at 61 and 65~\\k, with only one dip centered at $\\sim$ 63~\\k.  From Fig. \\ref{HCO_spectra}-a it can be noticed that the HCO$^+$ spectrum on position ($-$60, $-$20), that is at the position of the UCHII region G45.45, has also a dip at about 60~\\k~similarly to what it is observed in the CS J=7--6 line towards the same position.  The HCO$^+$ spectra also show evidence of emission near position (60, $-$40). This molecular feature is in positional coincidence with an infrared (IR) source located at RA=19$^{\\rm h}$14$^{\\rm m}$27.7$^{\\rm s}$, dec.=+11\\d08\\m33\\s (J2000) (see Fig \\ref{intro}). As this HCO$^+$ spectrum has the same kinematic structure than the spectrum towards ($-$60, $-$20) including the presence of a dip at the same velocity of about 60~\\k, we can conclude that this IR source is embedded in a molecular filament that must be connected with G45.45.  Figure \\ref{HCO-maps} shows the velocity channel maps of the HCO$^+$ J=4--3 emission averaged every 1~\\k. The HCO$^+$ J=4--3 emission related to G45.47 is visible between $\\sim$ 55 and $\\sim$ 68~\\k, while the emission associated with G45.45 goes from $\\sim$ 54 to $\\sim$ 63~\\k, respectively. Between $\\sim$ 56 and $\\sim$ 60~\\k part of the molecular condensation related to the IR source mentioned above can be appreciated.  \\begin{figure*} \\centering \\includegraphics[width=14cm]{HCO_spectra.eps} \\caption{a) HCO$^+$ J=4--3 spectra obtained towards the 2\\m~$\\times$~2\\m~region (white box in Fig. \\ref{intro}-left) mapped with ASTE. b) Hanning smoothed profile of the HCO$^+$ J=4--3 line towards the position (0, 0) where it is G45.47. The dashed line indicates the systemic velocity of the gas. The spectrum was deconstructed using three Gaussians, which are shown in red.} \\label{HCO_spectra} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=12cm]{HCO-maps.eps} \\caption{Velocity channel maps of the HCO$^+$ J=4--3 emission averaged every 1~\\k. Greyscale goes from 0.15 (about 3$\\sigma$ rms noise of an averaged channel map) to 2.5~K. Contours levels are at 0.15, 0.5, 0.8, 1.3, and 1.9~K. The red cross represents the position of G45.47.} \\label{HCO-maps} \\end{figure*}  Figure \\ref{13CO_spectra} shows the $^{13}$CO J=3--2 spectra towards the same region.  The $^{13}$CO emission towards (0, 0) can be decomposed in three velocity components centered at about 56 (the systemic velocity of the gas), 60, and 63~\\k, while only one dip centered at $\\sim$ 58~\\k can be appreciated. The depression at 58~\\k~can be observed in almost all the $^{13}$CO J=3--2 spectra in the region.  {\\bf Figure \\ref{13CO-maps} shows the velocity channel maps of the $^{13}$CO J=3--2 emission averaged every 1~\\k.}  \\begin{figure*} \\centering \\includegraphics[width=15cm]{13CO_spectra.eps} \\caption{$^{13}$CO J=3--2 spectra obtained towards the 2\\m~$\\times$~2\\m~region (white box in Fig. \\ref{intro}-left) mapped with ASTE. b) Hanning smoothed profile of the $^{13}$CO J=3--2 line towards G45.47.  The dashed line indicates the systemic velocity of the gas. The spectrum was deconstructed using three Gaussians, which are shown in red.} \\label{13CO_spectra} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=12cm]{13CO-maps.eps} \\caption{Velocity channel maps of the $^{13}$CO J=3--2 emission averaged every 1~\\k.  Greyscale goes from 0.15 up to 9~K. Contours levels are at 0.6, 2, 4, 5, 6, 7, 8, and 9~K. The red cross represents the position of G45.47.} \\label{13CO-maps} \\end{figure*}   The $^{12}$CO J=3--2 spectra (Fig. \\ref{12CO_spectra}) exhibit a more complicated behaviour. The profile towards the position (0, 0) has a triple peak structure with components centered at about 55, 60, and 64~\\k and two dips at about 58 and 62~\\k. This spectrum is broadened, suggesting the presence of outflowing activity in the region with the blue wing centered near the position (20, 20) and the red wing around the (0, $-$40) offset.  In Figure \\ref{12CO_maps} we show the velocity channel maps of the $^{12}$CO J=3--2 emission averaged every 2.1~\\k. Among all the observed molecular condensations, we draw the attention onto the clumps related to G45.47 and G45.45. The molecular clump related to G45.47 is observed between $\\sim$ 62 and 66~\\k and is seen slightly shifted to the northwest between $\\sim$ 58 and 62~\\k. The clump associated with G45.45 (partially observed) is centered at RA=19$^{\\rm h}$14$^{\\rm m}$22$^{\\rm s}$, dec.=+11\\d09\\m20\\s (J2000) in the velocity interval going from 47 to 63~\\k. Although both molecular condensations have associated different velocity ranges, they are connected through the extended emission shown in the velocity interval going from 50 to 65~\\k. The molecular gas related to the spectral wings appears as two conspicuous molecular features centered at RA=19$^{\\rm h}$14$^{\\rm m}$27$^{\\rm s}$, dec.=+11\\d09\\m45\\s (J2000) between 35 and 53~\\k~ and at RA=19$^{\\rm h}$14$^{\\rm m}$26$^{\\rm s}$, dec.=+11\\d08\\m45\\s (J2000) between 65 and 76~\\k.  These features will be further discussed in Section \\ref{outflows}.  \\begin{figure*} \\centering \\includegraphics[width=14cm]{12CO_spectra.eps} \\caption{a) $^{12}$CO J=3--2 spectra obtained towards the 2\\m~$\\times$~2\\m~region (white box in Fig. \\ref{intro}-left) mapped with ASTE. b) Hanning smoothed profile of the $^{12}$CO J=3--2 line towards G45.47. The dashed line indicates the systemic velocity of the gas. The spectrum was deconstructed using five Gaussians, which are shown in red.} \\label{12CO_spectra} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width=12cm]{12COJ=3-2-map.eps} \\caption{Velocity channel maps of the $^{12}$CO J=3--2 emission averaged every 2.1~\\k. The given velocities correspond to the higher velocity of each interval. Greyscale goes up to 23~K. Contours are above the 5$\\sigma$ of the rms noise level. The cross indicates the position of G45.47.} \\label{12CO_maps} \\end{figure*}  The four transitions have, within errors, the same main velocity components at about 55, 60, and 64~\\k~towards the position (0, 0). By the other hand, the velocity components observed at about 51 and 68~\\k~in the $^{12}$CO J=3--2 spectrum are not detected in the other three lines. Table \\ref{gaussianfit} lists the emission peaks parameters derived from a Gaussian fitting for the four molecular transitions on the position (0, 0).  T$_{mb}$ represents the peak brightness temperature and V$_{LSR}$ the central velocity referred to the Local Standard of Rest. Errors are formal 1$\\sigma$ value for the model of Gaussian line shape. All the spectra towards the (0,0) position have the same self-absorption dip at about 58-59~\\k, which correspond to the central velocity of the molecular cloud GRSMC 045.49+00.04 where G45.47 is embedded. Even more, this spectral feature is observed in all transitions towards the region G45.45 and the IR source located at RA=19$^{\\rm h}$14$^{\\rm m}$27.7$^{\\rm s}$, dec.=+11\\d08\\m33\\s (J2000) (see Fig \\ref{intro}).  \\begin{figure*} \\centering \\includegraphics[width=13cm]{wings_spectra.eps} \\caption{Comparison between the $^{12}$CO J=3--2 and the HCO$^+$ J=4--3 spectra on the positions (20, 20) and (0, -40). The dashed line indicates the systemic velocity. The velocity intervals corresponding to the blue and red wings are shown. The four spectra were Hanning smoothed.} \\label{spectralwings} \\end{figure*}   \\begin{table} \\caption{Emission peaks parameters derived from a Gaussian fitting for the four molecular transitions spectra on the position (0, 0).} \\centering \\begin{tabular}{cccc} \\hline Transition & V$_{LSR}$ [\\k] &  T$_{mb}$ [K] & $\\Delta$v [\\k]\\\\ \\hline CS J=7--6 & 56.3$\\pm$0.9 & 0.6$\\pm$0.2 & 3.1$\\pm$0.6\\\\ & 62.0$\\pm$0.8 & 1.8$\\pm$0.1 & 3.1$\\pm$0.7\\\\ & 64.8$\\pm$1.1 & 1.3$\\pm$0.1 & 2.5$\\pm$0.9\\\\ \\hline HCO$^+$ J=4--3  & 55.8$\\pm$1.2 & 0.4$\\pm$0.2 & 5.9$\\pm$0.5\\\\ & 60.8$\\pm$1.4 & 4.4$\\pm$0.8 & 3.8$\\pm$0.6\\\\ & 64.7$\\pm$1.7 & 5.1$\\pm$0.6 & 3.3$\\pm$0.7\\\\ \\hline $^{13}$CO J=3--2 & 55.9$\\pm$0.6 & 5.2$\\pm$0.4 & 2.9$\\pm$0.4\\\\ & 60.5$\\pm$0.5 & 11.1$\\pm$0.8 & 3.3$\\pm$0.6\\\\ & 63.3$\\pm$0.4 & 6.4$\\pm$0.7 & 2.6$\\pm$0.7\\\\ \\hline $^{12}$CO J=3--2 & 50.7$\\pm$0.3 & 3.0$\\pm$0.2 & 9.0$\\pm$0.3\\\\ & 55.2$\\pm$0.1 & 13.1$\\pm$1.1 & 3.3$\\pm$0.4\\\\ & 59.8$\\pm$0.2 & 10.6$\\pm$0.8 & 3.6$\\pm$0.6\\\\ & 64.4$\\pm$0.6 & 10.6$\\pm$0.7 & 3.3$\\pm$0.5\\\\ & 68.1$\\pm$0.4 & 1.3$\\pm$0.2 & 4.4$\\pm$0.4\\\\ \\hline  \\label{gaussianfit} \\end{tabular} \\end{table}     \\subsection{Column density and mass estimates of the molecular clump associated with G45.47} \\label{clump}  We estimate the $^{13}$CO J=3--2 opacity, $\\tau_{13}$, based on the following equation:  \\begin{equation} \\small \\tau_{13}=-ln\\left(1-\\frac{T_{peak}(^{13}{\\rm CO})}{T_{peak}(^{12}{\\rm CO})}\\right) \\end{equation}  \\noindent where we consider T$_{peak}$ from the position (0, 0). We obtain, $\\tau_{13} \\sim$ 1.9 which reveals that the $^{13}$CO J=3--2 emission is optically thick towards G45.47, in agreement with the observed profile towards (0,0) offset (see Fig. \\ref{13CO_spectra}).  The excitation temperature, $T_{ex}$, of the $^{13}$CO J=3--2 line was estimated from:  \\begin{equation} \\small ^{13}T_{peak}=\\frac{h\\nu}{k} \\left(\\frac{1}{e^{h\\nu/k T_{ex}}-1}-\\frac{1}{e^{h\\nu/k T_{BG}}-1}\\right) \\times (1-e^{-\\tau_{13}}) \\end{equation}  \\noindent where for this line $h\\nu/k=15.87$. Assuming $T_{BG}$ = 2.7~K, and considering the peak brightness temperature for the $^{13}$CO J=3--2 at (0, 0) offset, $^{13}T_{peak}$ = 11.15~K, we derive a $T_{ex} \\sim$ 20~K for the $^{13}$CO J=3--2 line.  Finally, using the RADEX \\footnote{http://www.sron.rug.nl/$\\sim$vdtak/radex/radex.php} code \\citep{tak07} we derive the $^{13}$CO J=3--2 column density and the H$_2$ volume density. The RADEX model uses the mean escape probability approximation for the radiative transfer equation.  Adopting $T_{ex} \\sim$ 20~K, $\\tau_{13} \\sim$ 1.9, and $^{13}T_{peak}$ = 11.15~K, we obtain a $^{13}$CO column density N($^{13}$CO) $\\sim 2.8 \\times 10^{17} {\\rm cm}^{-2}$ and n(H$_2$) $\\sim 10^5 {\\rm cm}^{-3}$. Considering an abundance ratio of [H$_2$]/[$^{13}$CO] = 77$\\times 10^4$ \\citep{wil94} we estimate the H$_2$ column density, N(H$_2) \\sim 2.1 \\times 10^{23}{\\rm cm}^{-2}$. Finally, using the relation $M=\\mu m_Hd^2\\Omega {\\rm N(H_2)}$, where $\\mu$ is the mean molecular weight per H$_2$ molecule ($\\mu \\sim 2.72$), $m_H$ the hydrogen atomic mass, $d$ the distance, and $\\Omega$ the solid angle subtended by the structure, we calculate the total mass of the clump in $M \\sim 10^4$ \\msol.  As an independent estimate, we derive the beam-averaged gas column density, the mass, and the volume density of the clump based on the dust continuum emission. In particular, we use the integrated flux of the continuum emission at 1.1~mm as obtained from The Bolocam Galactic Plane Survey II Catalog (BGPS II; \\citealt{ros10}). The 1.1~mm continuum emission is mostly originated in optically thin dust \\citep{hild83}. Following \\citet{be02} and \\citet{hild83} we calculate the mass and the gas column density of the clump using:  \\begin{eqnarray} M_{gas}=\\frac{1.3 \\times 10^{-3}}{J_{\\nu}(T_{dust})}\\frac{a}{0.1 \\mu{\\rm m}}\\frac{\\rho}{3 {\\rm g cm}^{-3}}\\frac{R}{100}\\frac{F_{\\nu}}{{\\rm Jy}}\\nonumber\\\\\\left(\\frac{d}{{\\rm kpc}}\\right)^2 \\left(\\frac{\\nu}{2.4 {\\rm THz}}\\right)^{-3-\\beta} [{\\rm M}_{\\odot}] \\end{eqnarray}  \\noindent and  \\begin{eqnarray} N_{gas}=\\frac{7.8 \\times 10^{10}}{J_{\\nu}(T_{dust})\\Omega_b}\\frac{a}{0.1 \\mu{\\rm m}}\\frac{\\rho}{3 {\\rm g cm}^{-3}}\\frac{R}{100}\\frac{F_{\\nu}}{{\\rm Jy}}\\nonumber\\\\ \\left(\\frac{\\nu}{2.4 {\\rm THz}}\\right)^{-3-\\beta} [{\\rm cm}^{-2}] \\end{eqnarray}   \\noindent where $J_{\\nu}(T_{dust}$) = [exp($h\\nu/kT_{dust})-1]^{-1}$ and $\\Omega_b, a, \\rho, R,$ and $\\beta$ are the beam solid angle, grain size, grain mass density, gas-to-dust ratio, and grain emissivity index for which we used the values of (33\\s)$^2$ in radians, 0.1~$\\mu$m, 3~g cm$^{-3}$, 100, and 2, respectively (\\citealt{hun97}, \\citealt{hun00}, and \\citealt{mol00}). Based on the work of \\citet{sri02} who derived dust temperatures ranging between 30 and 60~K for a sample of several massive star forming regions, we adopt $T_{dust}$ = 45~K. For a distance $d$ = 8.3~kpc and an integrated flux intensity $F_{\\nu}$ = 5.18~Jy at 1~mm \\citep{ros10} we obtain $N_{gas} \\sim 4 \\times 10^{23}$~cm$^{-2}$, $M_{gas} \\sim 8520$~\\msol, and a volume density, $n$(H$_2) \\sim 1.4 \\times 10^5$~cm$^{-3}$.  These values are in good agreement with those derived from the $^{13}$CO J=3--2 line using RADEX.    \\subsection{Molecular outflows associated with G45.47} \\label{outflows}  As discussed in Section \\ref{molecular}, the presence of spectral wings in the $^{12}$CO J=3--2 spectrum obtained towards G45.47 is a strong indicator that molecular outflow activity is taking place in the region. To characterize the associated molecular outflows it is first necessary to separate the outflowing gas from the molecular material of the clump, identifying the velocity ranges related to each structure.  We consider two independent methods to determine these velocity ranges. The first method consists in the identification of the blue and red spectral wings based on a comparison between the $^{12}$CO J=3--2 and HCO$^+$ J=4--3 spectra. Figure \\ref{spectralwings} shows the comparison between both spectra on positions (20, 20) and (0, $-$40) where the blue and red wings are largest.  Considering emission up to about 2~$\\sigma$ of rms noise level, it can be noticed a blue and a red wing in the $^{12}$CO spectra between $\\sim$ 34 and $\\sim$ 54 and between $\\sim$ 64 and $\\sim$ 75~\\k, respectively.    The second method to identify the molecular outflows is based on the inspection of the $^{12}$CO J=3--2 data cube, channel by channel, trying to spatially separate both outflows's lobes. In Section \\ref{molecular} we mentioned two molecular features in the $^{12}$CO J=3--2 emission centered at RA=19$^{\\rm h}$14$^{\\rm m}$27$^{\\rm s}$, dec.=+11\\d09\\m45\\s (J2000) between $\\sim$ 34 and $\\sim$ 54~\\k~and at RA=19$^{\\rm h}$14$^{\\rm m}$26$^{\\rm s}$, dec.=+11\\d08\\m45\\s (J2000) between $\\sim$ 64 and $\\sim$ 75~\\k~which we identify as the outflow's lobes related to G45.47 (see Fig. \\ref{12CO_maps}). These molecular features can be identified as the spectral wings detected in the $^{12}$CO J=3--2 spectra through the first method.  Figure \\ref{blue-red-wings} shows a {\\it Spitzer}-IRAC three color image (3.6~$\\mu$m = blue, 4.5~$\\mu$m = green and 8~$\\mu$m = red) of G45.47. The blue and red contours represent the $^{12}$CO J=3--2 emission integrated from $\\sim$ 34 to 54~\\k~(blue lobe) and from $\\sim$ 64 to 75~\\k~(red lobe), respectively.   \\begin{figure}[h] \\centering \\includegraphics[width=8cm]{blue-red-wings2.eps} \\caption{{\\it Spitzer}-IRAC three color image (3.6~$\\mu$m = blue, 4.5~$\\mu$m = green and 8~$\\mu$m = red) of G45.47. The blue and red contours represent the $^{12}$CO J=3--2 emission integrated from 34 to 54~\\k~(blue lobe), and from 64 to 75~\\k~(red lobe), respectively. The blue contours are at 400, 500, 600, and 670 K~\\k~ and the red ones are at 140, 160, 180, and 200 K~\\k.} \\label{blue-red-wings} \\end{figure}   Following \\citet{sco86} we estimate the $^{12}$CO optical depth, $\\tau_{12}$, of the gas in the molecular outflows using:  \\begin{equation} \\small \\frac{^{12}{\\rm T}_{mb}}{^{13}{\\rm T}_{mb}}=\\frac{1-exp(-\\tau_{12})}{1-exp(-\\tau_{12}/X)} \\end{equation}   \\noindent where $X$=[$^{12}$CO]/[$^{13}$CO] is the isotope abundance ratio. Using the relation $X=6.21 \\times D_{GC} + 18.71$ \\citep{mil05}, where D$_{GC}$ = 6.1~kpc is the distance between G45.47 and the Galactic Center, we obtain an isotope abundance ratio $X$= 57 for this region. We adopt a constant $X$ throughout the outflows \\citep{cab90}. Since we could not identify any blue wing in the $^{13}$CO spectrum on position (20, 20), the $^{12}$T$_{mb}/^{13}$T$_{mb}$ ratio for both wings was estimated considering only the red one. We obtain a $^{12}$T$_{mb}/^{13}$T$_{mb}$ ratio of about 4.2 and derive a $^{12}$CO optical depth of $\\tau_{12} \\sim$ 15, which will be adopted for both spectral wings in further calculations.  We can now estimate the $^{12}$CO column density of both outflow lobes from (see e.g. \\citealt{buc10}):  \\begin{equation} \\small {\\rm N}(^{12}{\\rm CO})=7.96\\times10^{13} e^{\\frac{16.6}{T_{ex}}}\\frac{T_{ex}+0.92}{1-exp(\\frac{-16.6}{T_{ex}})} \\int \\tau_{12} {\\rm dv} \\end{equation}  \\noindent Taking into account that the $^{12}$CO J=3--2 is an optically thick transition ($\\tau \\geq 1$), the integral can be approximated by:  \\begin{equation} \\small \\int \\tau_{12} {\\rm dv} = \\frac{1}{J(T_{ex})-J(T_{BG})}\\frac{\\tau_{12}}{1-e^{-\\tau_{12}}}\\int {\\rm ^{12}T}_{mb}~{\\rm dv} \\label{thick} \\end{equation}  \\noindent By integrating the $^{12}$CO emission from 34 to 54~\\k~and from 64 to 75~\\k~we obtain N(CO)$_{blue} \\sim 5.2 \\times 10^{17} {\\rm cm}^{-2}$ and N(CO)$_{red} \\sim 2 \\times 10^{17} {\\rm cm}^{-2}$, respectively. The mass of each wing can be derived from the relation $M=\\mu m_Hd^2\\Omega X({\\rm CO})^{-1}{\\rm N(CO)}$, where $X$(CO) is the $^{12}$CO relative abundance to H$_2$ ($X$(CO)$\\sim 7.4 \\times 10^{-5}$). We obtain M$_{blue} \\sim 300$ \\msol~and M$_{red} \\sim 120$ \\msol, yielding a total outflow mass of about 420~\\msol. In this way, the estimated total outflows mass represents about 4\\% of the clump mass, which is in agreement with the results of \\citet{beutPhD02} who established that approximately 4\\% of the clump gas is entrained in the molecular outflows.  The momentum and the kinetic energy of the wings can be derived from $P = M V_c$ and $E_k = 0.5 M V_c^2$, where $V_c$ is a characteristic velocity estimated as the difference between the maximum velocity of detectable $^{12}$CO emission in the wing and the systemic velocity of the gas ($\\sim$ 56~\\k). Taking into account a $V_c^{blue} \\sim$ 22~\\k~ and a $V_c^{red} \\sim$ 19~\\k, we obtain $P_{blue} = 6.6 \\times 10^3$ \\msol \\k, $P_{red} = 2.3 \\times 10^3$ \\msol \\k, $E_k^{blue} = 1.4 \\times 10^{48}$ erg and $E_k^{red} = 4.2 \\times 10^{47}$ erg. Comparing our results with the work of \\citet{wu04} who based on compiled data concluded that the mass, momentum, and energy of molecular outflows range from 10$^{-3}$ to 10$^{3}$~\\msol, 10$^{-3}$ to 10$^{4}$ \\msol \\k~and 10$^{38}$ to 10$^{48}$ erg, respectively, we can infer that in this work we are dealing with massive and energetic outflows.  Following \\citet{curtis10} we define the dynamical time for the blue and red wings, $t_{dyn}$, as the time for the bow shock travelling at the maximum velocity in the flow, $V_c$, to travel the projected lobe length, $R_{lobe}$:  \\begin{equation} t_{dyn}=\\frac{R_{lobe}}{V_c} \\end{equation}  \\noindent with $R_{lobe}^{blue}$ = 1.8~pc and $R_{lobe}^{red}$ = 1.6~pc. These values were obtained by inspecting the blue and red lobe sizes in Fig. \\ref{blue-red-wings} and considering a distance of 8.3~kpc. We obtain a similar dynamical time of both lobes, $t_{dyn} = 0.8 \\times 10^5$~yr. According to \\citet{beuther02}, flow ages are good estimates of protostar lifetimes. We can then assume that the protostar responsible of the outflow activity has an age of the order of $10^5$~yr. This is consistent with the detection of ionized gas in the region, since to be able to ionize its surroundings the age of the protostar should be at least $10^5$~yr \\citep{sri02}.  \\subsection{HCO$^{+}$ as tracer of molecular outflows} \\label{hco+}  As an independent test of out findings, we performed a study of the HCO$^+$ abundance towards G45.47. HCO$^{+}$ is believed to be the dominant ionized species in dense dark clouds \\citep{rawlings00,dalgarno84}, being very important for the ion-neutral chemistry. The ionization fraction of molecular clouds is a relevant parameter to study the cloud chemistry and dynamics, and hence the star forming processes.  Star formation occurs in the interior of dense cores, regions of high extinction where self-shielding prevents the UV photoionization of H$_{2}$. Thus, cosmic ray ionization is believed to dominate photoionization in dense cores \\citep{mckee89}.  However, when YSOs are formed within the cores, the shocks and outflows produce deep changes in the chemistry and photoionization processes are very likely.  Indeed, in star forming regions occur a molecular enrichment due to the desorption of molecular-rich ice mantles, followed by photochemical processing by shock-generated radiation fields.  In particular, towards YSOs which are driving outflows, it is expectable an enhancement in the HCO$^{+}$ abundance \\citep{rawlings00,rawlings04}. In what follows, we estimate, using a simple chemical network, the HCO$^{+}$ abundance that would be produced by an standard cosmic ray ionization rate in the G45 clump in order to compare with the abundance obtained from our observations.  The starting chemical reaction to form HCO$^{+}$ is the production of H$_{3}^{+}$, mainly formed in the interaction between the cosmic rays (c.r.) and the molecular gas (e.g. \\citealt{Oka06}): \\begin{eqnarray} \\centering {\\rm H_{2} + c.r.}  &\\rightarrow&  {\\rm H_{2}^{+} + e^{-} }  \\nonumber \\\\ {\\rm H_{2}^{+} + H_{2} }   &\\rightarrow&  {\\rm H_{3}^{+} + H. } \\label{chem1} \\end{eqnarray} The rate equation of this reaction is \\begin{equation} \\zeta_{\\rm H_{2}}~n({\\rm H_{2}}) = \\frac{dn({\\rm H_{3}^{+}})}{dt}, \\label{rate1} \\end{equation} where $\\zeta_{\\rm H_{2}}$ is the cosmic ray ionization rate and $n$ the density of the molecular species. By considering that the HCO$^{+}$ is mainly formed by the reaction of H$_{3}^{+}$ with CO: \\begin{equation} {\\rm H_{3}^{+} +  CO \\rightarrow  HCO^{+} + H_{2} } \\label{chem2} \\end{equation} and destroyed through recombination with electrons, \\begin{equation} {\\rm HCO^{+} + e^{-} \\rightarrow CO + H,  } \\label{chem3} \\end{equation} the rate equations can be equated, leading to \\begin{equation} k_{\\rm HCO^{+}}~n({\\rm H_{3}^{+}})~n({\\rm CO}) = k_{\\rm CO}~n({\\rm HCO^{+}})~n(e), \\label{rate2} \\end{equation} where $k_{\\rm HCO^{+}}$ and $k_{\\rm CO}$ are the coefficient rates and $n(e)$ the electron density. Assuming that the main destruction mechanism for H$_{3}^{+}$ is through the formation of HCO$^{+}$, the rate of destruction of H$_{3}^{+}$ is therefore equal to the formation rate of HCO$^{+}$; this implies that the left side of equation (\\ref{rate1}) is equal to right side of equation (\\ref{rate2}). Then, it is possible to write an expression for the cosmic ionization rate as a function of the molecular densities: \\begin{equation} \\zeta_{\\rm H_{2}} =  k_{\\rm CO} \\frac{n({\\rm HCO^{+}})~n(e)}{n({\\rm H_{2}})}. \\label{chem4} \\end{equation} The rate coefficient $k_{\\rm CO}$, extracted from the UMIST database \\citep{woodall07}, is: $k_{\\rm CO} = 2.4 \\times 10^{-7} (T/300 {\\rm K})^{-0.69}$. Equation (\\ref{chem4}) can be approximated using the column densities, leading: \\begin{equation} \\zeta_{\\rm H_{2}} \\simeq  k_{\\rm CO} \\frac{{\\rm N(HCO^{+})}~X(e)~n({\\rm H_{2}})}{\\rm N(H_{2})}, \\label{chem5} \\end{equation} where $X(e)$ is the electron abundance. Using this equation and assuming typical values for the cosmic ionization rate and electron abundance of $\\zeta_{\\rm H_{2}} = (1-5) \\times 10^{-17}$ s$^{-1}$ \\citep{dalgarno06} and $X(e) \\sim 10^{-7}$ \\citep{bergin99}, respectively, we derive the HCO$^{+}$ column density. We use N(H$_{2}$) $= 2 \\times 10^{23}$ cm$^{-2}$ and $n({\\rm H_{2}}) = 1 \\times 10^{5}$ cm$^{-3}$, as derived above, and we assume $T = 20$ K. Finally, we obtain that the N(HCO$^{+}$) should be in the range $(1.5 - 6.5) \\times 10^{14}$ cm$^{-2}$, leading to an abundance $X({\\rm HCO^{+}})$ in between $6.5 \\times 10^{-10}$ and $3.2 \\times 10^{-9}$.  On the other hand, we analyze the central HCO$^{+}$ J=4--3 spectrum and use the RADEX code \\citep{tak07} to derive its column density. Assuming background and kinetic temperatures of 2.73 K and 20 K, respectively, and the same H$_{2}$ density used above, we varied the HCO$^{+}$ column density until obtaining a good fit for the observed peak temperature. The best fit was obtained with N(HCO$^{+}$)$ \\sim 1.5 \\times 10^{15}$ cm$^{-2}$, leading an abundance of $X({\\rm HCO^{+}}) \\sim 7.5 \\times 10^{-9}$.  This value, more than twice greater than the lager value obtained above, might indicate that the cosmic ray ionization is insufficient to produce the observed HCO$^{+}$ abundance through the described chemical network. Therefore, following \\citet{rawlings00,rawlings04}, we suggest that the observed HCO$^{+}$ abundance must be mostly produced by the outflowing activity in the region. In spite of the uncertainties in the calculations, this is an independent proof pointing to support an scenario with outflows in the region.   \\subsection{Looking for the driving source of the molecular outflows}  \\citet{wil96} reported 5\\s~resolution observations of the HCO$^+$ J=1--0 transition towards G45.47 obtained with the OVRO millimeter array. They identified at least five HCO$^+$ J=1--0 clumps and suggested that the formation of an OB cluster would be taking place in the region.  However the authors did not carry out any search of infrared point sources in the region to confirm that hypothesis.  In this context, we wonder if the massive molecular outflows observed towards G45.47 were originated by one or more stars. Based on a near-infrared analysis, we searched for the driving source candidates of the massive molecular outflows.    Figure \\ref{HCO-clumps} shows the radio continuum emission at 6~cm obtained from The Multi-Array Galactic Plane Imaging Survey (MAGPIS; \\citealt{white05}) of the area mapped using ASTE. The green contours represent the HCO$^+$ J=4--3 emission distribution integrated between 58 and 67~\\k. The black contours were taken from the paper of \\citet{wil96} and represent the HCO$^+$ J=1--0 emission distribution integrated in the same velocity interval.  The region observed by \\citet{wil96} is indicated by the dashed box.  The five clumps reported by the authors have been labeled.  The positional coincidence among the center of the HCO$^+$ J=4--3 clump, the clump 3 of HCO$^+$ J=1--0 and a radio continuum source, which has been identified as the UCHII region G45.47, is striking.  The red circles indicate the location of the YSO candidates in the observed region (see red circles in Fig. \\ref{2mass}). Among the five YSO candidates projected onto the HCO$^+$ J=4--3 emission, 2MASS 19142564+1109283 is the only one having a positional coincidence with a clump of HCO$^+$ J=1--0 (clump 3). We do not find any embedded infrared sources in the others four clumps, suggesting that these clumps could be in a prestellar core stage or they might be tracing the origin of the molecular outflows. We suggest that the most likely candidate to be the driving source of the molecular outflows is 2MASS 19142564+1109283.  \\begin{figure}[h] \\centering \\includegraphics[width=7.5cm]{HCO-wilner.eps} \\caption{Radio continuum emission at 6~cm (MAGPIS; \\citealt{white05}) of the area mapped using ASTE. The green contours represent the HCO$^+$ J=4--3 emission distribution integrated between 58 and 67~\\k. The black contours were taken from the paper of \\citet{wil96} and represent the HCO$^+$ J=1--0 emission distribution integrated from 58 and 67~\\k. The box indicates the area observed by \\citet{wil96}. The numbers show the five HCO$^+$ J=1--0 clumps. The red circles represent the YSO candidates shown in Fig. \\ref{2mass}.} \\label{HCO-clumps} \\end{figure}  Figure \\ref{2mass} shows a color-color diagram (CCD) including all the 2MASS point sources located in the observed region. The red circles represent the sources with infrared excess (YSO candidates) which have been shown in Figure \\ref{HCO-clumps}. The reddenest object is 2MASS 19142564+1109283. The blue circles represent main sequence or giant star candidates.  \\begin{figure}[h] \\centering \\includegraphics[width=7.5cm, angle=-90]{2MASS_CC.ps} \\caption{Color-color diagram of the 2MASS infrared sources in the vicinity of G45.47. The two solid curves represent the location of the main sequence (thin line) and the giant stars (thick line) derived from \\citet{bessell88}. The parallel dashed lines are reddening vectors with the crosses placed at intervals corresponding to five magnitudes of visual extinction. We have assumed the interstellar reddening law of \\citet{rieke85} (A$_J$/A$_V$ =0.282; A$_H$/A$_V$ =0.175 and A$_K$/A$_V$ =0.112).  The sources reddened by circumstellar dust are indicated as red circles.} \\label{2mass} \\end{figure}     To characterize the infrared source 2MASS 19142564+1109283 we perform a fitting of its spectral energy distribution (SED) using the tool developed by \\citet{rob07}\\footnote{http://caravan.astro.wisc.edu/protostars/}. We adopt an interstellar extinction in the line of sight, A$_v$, between 5 and 17 magnitudes. The range of A$_v$ was chosen by inspecting the location in a [{\\it H-K}] vs [{\\it J-H}] CCD of the 2MASS sources in the region (see Fig. \\ref{2mass}).  To construct the SED we consider the fluxes at the {\\it JHK} 2MASS bands, {\\it Spitzer}-IRAC 5.8 and 8 $\\mu$m bands, WISE 3.4, 4.6, 12 and 22 $\\mu$m bands, MSX 14 and 21 $\\mu$m bands, and SCUBA 850 $\\mu$m band.  The fluxes of the datasets having lower angular resolution (MSX and SCUBA) were considered as upper limits.  Figure \\ref{SED} shows the best fitting SEDs models for 2MASS 19142564+1109283. We select models that satisfies the condition:  \\begin{equation} \\chi^2-\\chi_{best}^2 < 2N, \\label{selmol} \\end{equation}  \\noindent where $\\chi_{best}^2$ is the minimum value of the $\\chi^2$ among all models, and $N$ is the number of input data fluxes (fluxes specified as upper limit do not contribute to $N$). Hereafter, we refer to models satisfying Eq. \\ref{selmol} as ``selected models''.  \\begin{figure}[h] \\centering \\includegraphics[width=8cm]{SED.eps} \\caption{Best fitting SED models for G45.47. The fill circles indicate the measured fluxes at the {\\it JHK} 2MASS bands, {\\it Spitzer}-IRAC bands at 5.8 and 8.0 $\\mu$m, and the WISE bands at 3.4, 4.6, 12 and 22 $\\mu$m. Triangles indicate fluxes considered as upper limits, as MSX bands at 14 and 21 $\\mu$m and the SCUBA band at 850 $\\mu$m. Black and gray solid curves represent the best-fit model and the subsequent good fittings, respectively. The dashed line shows the best-fit model for a central source contribution in absence of circumstellar dust.} \\label{SED} \\end{figure}   From the SED fitting several model parameters can be obtained and, as expected, some can be better constrained than others.  Taking into account the values yielding by the ``selected models'' for several parameters, we derive the following main results:  \\begin{itemize}  \\item the total luminosity distribution of the source has an average value of $2 \\times 10^4$~\\lsol~with a spread between $8 \\times 10^3$ and $5 \\times 10^4$~\\lsol.  \\item embedded in the molecular clump there is a massive protostar of about 15 \\msol~(early B-type star) accreting material from its envelope at large rates, $\\dot M_{env} \\sim 2\\times10^{-4}$~\\msol~yr$^{-1}$.  \\item despite the large spread in the central source's age distribution that goes from $2 \\times 10^3$ to $10^6$ yr, it peaks at about $3\\times 10^5$ yr, as expected for protostars that has begun to ionize their surroundings. Besides, this result is in agreement with the dynamical time, $t_{dyn} = 1 \\times 10^5$~yr, derived in Section \\ref{outflows}.  \\end{itemize}"
    },
    "1208/1208.5987_arXiv.txt": {
        "abstract": "We present Gemini/GNIRS spectroscopic observations of 4 $z-$band ($z\\approx 7$) dropout galaxies and VLT/XSHOOTER observations of one $z-$band dropout and 3 $Y-$band ($z\\approx 8-9$) dropout galaxies in the {\\em Hubble} Ultra Deep Field, which were selected with Wide Field Camera 3 imaging on the {\\em Hubble Space Telescope}. We find no evidence of Lyman-$\\alpha$ emission with a typical $5\\sigma$ sensitivity of $5\\times10^{-18}$erg cm$^{-2}$s$^{-1}$, and we use the upper limits on Lyman-$\\alpha$ flux and the broad-band magnitudes to constrain the rest-frame equivalent widths for this line emission. Accounting for incomplete spectral coverage, we survey 3.0 $z$-band dropouts and 2.9 $Y$-band dropouts to a Lyman-$\\alpha$ rest-frame equivalent width limit $>120$\\,\\AA\\ (for an unresolved emission line); for an equivalent width limit of $50$\\,\\AA\\ the effective numbers of drop-outs surveyed fall to 1.2 $z$-band drop-outs and 1.5 $Y$-band drop-outs.  A simple model where the fraction of high rest-frame equivalent width emitters follows the trend seen at $z=3-6.5$ is inconsistent with our non-detections at $z=7-9$   at the $\\approx 1\\sigma$ level for spectrally unresolved lines, which may indicate that a significant neutral HI fraction in the intergalactic medium suppresses the Lyman-$\\alpha$ line in $z$-drop and $Y$-drop galaxies at $z>7$. ",
        "introduction": "The Lyman break technique has proven to be an efficient tool in the selection of high-redshift candidate galaxies - and rest-frame UV-selected galaxies are now being regularly identified at redshifts $z\\gtrsim6$ (e.g. Bunker et al.\\ 2004; Bouwens et al.\\ 2006; McLure et al.\\ 2010).  With the advent of the Wide Field Camera 3 (WFC3) on the {\\em Hubble} Space Telescope ({\\em HST}), the technique has been pushed to even higher redshifts and the number of candidates beyond $z=6.5$ has increased to about a hundred (e.g. Bouwens et al.\\ 2011, McLure et al.\\ 2011, Lorenzoni et al.\\ 2011, Wilkins et al.\\ 2011, Finkelstein et al.\\ 2010, Wilkins et al.\\ 2010, Bunker et al.\\ 2010, Oesch et al.\\ 2010a).  Spectroscopic confirmation of these high-$z$ candidate galaxies remains a very important task.  However, the sample of convincing spectroscopically-confirmed sources remains small (eg. Ono et al.\\ 2012, Pentericci et al.\\ 2011, Schenker et al.\\ 2012, Vanzella et al.\\ 2011).  Not knowing the precise redshift of most Lyman break candidates by means of spectroscopic confirmation reduces our confidence in inferred properties about the universe at this epoch, both due to redshift uncertainties on each source and because of the unquantifiable risk of contamination in these samples (e.g. see Hayes et al.\\ 2012).  For objects at such high redshifts, the only currently-feasible spectroscopic redshift diagnostic is the Lyman-$\\alpha$ emission line, which arises from photoionization of HII regions by star formation.  The line itself is resonant and sensitive to the ionization state of the intergalactic medium (IGM) and thus, its visibility at the end of cosmic reionization is expected to be reduced compared to that at lower redshifts due to the damping effect of an increasingly neutral IGM (Gunn \\& Peterson 1965; Becker et al.\\ 2001).  The emergence or non-emergence of Lyman-$\\alpha$ may be indicative of the size of the HII regions surrounding these galaxies; in order for Lyman-$\\alpha$ emission to emerge from a galaxy, the ionized region surrounding the galaxy needs to be sufficiently large to allow the wavelength of the Lyman-$\\alpha$ photon to redshift to longer wavelengths, and hence become non-resonant when encountering the neutral IGM, thus managing to escape.  In order to address the question of the emergence of Lyman-$\\alpha$ at these high redshifts, we undertook Gemini/GNIRS and VLT/XSHOOTER spectroscopy of a sample of 7 Lyman-break selected candidate galaxies at $z\\gtrsim7$, identified as $z$-band and $Y$-band dropouts with {\\em HST}/WFC3 imaging (Bunker et al.\\ 2010, Lorenzoni et al.\\ 2011, Wilkins et al.\\ 2011) centered on and around the Hubble Ultra Deep Field (HUDF, Beckwith et al.\\ 2006).  We have already described the analysis of one of these objects (HUDF.YD3) in a separate paper (Bunker et al.\\ 2012), following the reported detection of Lyman-$\\alpha$ at $z=8.55$ by Lehnert et al.\\ (2010), so we do not discuss it again here.  The structure of this paper is as follows. We describe our spectroscopic observations and data reduction in Section \\ref{sec:obs}, and present the results of the spectroscopy in Section \\ref{sec:results}.  We discuss our analysis in Section \\ref{sec:analysis} and our conclusions are summarized in Section \\ref{sec:conc}.  We adopt a $\\Lambda$CDM cosmology throughout, with $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_0=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$. All magnitudes are in the AB system (Oke \\& Gunn 1983). ",
        "conclusions": "We have presented spectroscopic observations with GEMINI/GNIRS and VLT/XSHOOTER of a sample of $z > 7$ candidate galaxies and we fail to detect significant Lyman-$\\alpha$ emission from any of them.  This is consistent with the fraction of high rest-frame equivalent width Lyman-$\\alpha$ emitters dropping at $z>7$, as would be expected if the neutral HI fraction was greater at these epochs.  We have also investigated a tentative emission line published by Fontana et al.\\ (2010) in HUDF.zD1 (from the catalogue of Bunker et al.\\ 2010) and our analysis does not confirm the presence of this line, although we do not rule out the possibility of it being real, especially if it has considerable velocity extent.  Given the lack of Lyman-$\\alpha$ emission in our spectroscopy in conjunction with the continuum flux derived from {\\em HST} imaging of these objects, we derived upper limits on the rest-frame equivalent width of our objects.  Extrapolating the Lyman-$\\alpha$ fraction observed at lower redshifts by Stark et al.\\ (2010) and Shapley et al.\\ (2003), our lack of Lyman-$\\alpha$ detection rules out at a level of 1\\,$\\sigma$ (70\\%), for spectrally unresolved lines, the scenario in which the Lyman-$\\alpha$ fraction evolves with the same trend found at lower redshifts.  The limits are weaker if the lines have significant velocity width extent.  A diminished Lyman-$\\alpha$ fraction at higher redshift is consistent with other published studies.  This attenuation in the Lyman-$\\alpha$ fraction can be attributed either to physical evolution of the galaxies or, more likely, an increase in the neutral fraction of hydrogen at $z > 7$, i.e. these observations can most likely be interpreted as implying that the neutral fraction at $z\\sim8$ can be ruled out as being $\\chi_{\\scriptscriptstyle HI}=0$ at a level of 1\\,$\\sigma$.  Larger-number statistics are required to confirm this hypothesis at a higher level of significance.  To this end, we have undertaken spectroscopy on a large sample of $z$-band and $Y$-band dropouts with VLT/FORS2 and SUBARU/MOIRCS (Caruana et al.\\ {\\em in prep}).  \\label{sec:conc}  \\subsection*{Acknowledgements} Based on observations obtained at the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under a cooperative agreement with the NSF on behalf of the Gemini partnership: the National Science Foundation (United States), the Science and Technology Facilities Council (United Kingdom), the National Research Council (Canada), CONICYT (Chile), the Australian Research Council (Australia), Minist\\'{e}rio da Ci\\^{e}ncia, Tecnologia e Inova\\c{c}\\~{a}o (Brazil) and Ministerio de Ciencia, Tecnolog\\'{i}a e Innovaci\\'{o}n Productiva (Argentina).  The Gemini programmes associated with these observations were:  GS-2004B-Q-19 and GS-2005B-Q-18.  Based on observations made with the NASA/ESA {\\em Hubble} Space Telescope associated with programmes \\#GO-11563, \\#GO/DP-10086, \\#GO-9803, \\#GO-9425. obtained from the Data Archive at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555.  We gratefully acknowledge Chris Willott, Richard Ellis, Daniel Stark and the NIRSpec Instrument Science Team for useful discussions.  We thank the anonymous referee whose helpful and insightful comments greatly improved this manuscript.  MJJ acknowledges the support of a RCUK fellowship. JC and SL are supported by the Marie Curie Initial Training Network ELIXIR of the European Commission under contract PITN-GA-2008-214227. AB and SW acknowledge financial support from an STFC Standard Grant.  \\begin{table*} \\caption{This table for all objects targeted by our spectroscopy shows the redshift range spanned by our data for Lyman-$\\alpha$ in Column 2.  Column 3 shows the fractional probability that a galaxy drawn from the dropout sample would fall within the spectral coverage of that particular spectrograph setup.  Column 4 gives the median EW for each object for the most probable redshift range (see Figure \\ref{fig:sims}).  The remaining three columns tabulate the fraction of our spectroscopy that has EW limits lower than our chosen threshold (50\\,\\AA, 75\\,\\AA\\, and 120\\,\\AA\\, respectively.)  The figures in this table are for an unresolved line.} \\begin{tabular}{ | c | c | c | c | c | c | c | } \\hline Object & z-Range spanned & Frac$_{z}$ & Median EW & Frac$_{\\mathrm{EW < 50\\AA}}$ &  Frac$_{\\mathrm{EW < 75\\AA}}$ & Frac$_{\\mathrm{EW < 120\\AA}}$\\\\ % & by data (for Ly-$\\alpha$)& & & & & \\\\ \\hline \\hline zD1 \t\t\t& 5.82 -- 6.40 & 0.01 & 54.69 & 0 & 0.4583 & 1.0000 \\\\ & 6.08 -- 6.74 & 0.24 & 243.2 & 0 & 0 & 0 \\\\ & 6.79 -- 7.45 & 0.50 & 29.35 & 0.8349 & 0.9415 & 0.9990 \\\\ & 7.25 -- 8.05 & 0.36 & 85.29 & 0 & 0.3326 & 0.6299 \\\\ \\hline zD2\t\t\t& 6.09 -- 6.73 & 0.24 & 113.5 & 0 & 0 & 0.7697 \\\\ & 7.25 -- 8.00 & 0.33 & 76.10 & 0.285052 & 0.4109 & 0.6883 \\\\ \\hline zD3\t\t\t& 6.09 -- 6.73 & 0.24 & 120.4 & 0 & 0 & 0.7262 \\\\ & 7.25 -- 8.00 & 0.33 & 80.72 & 0 & 0.3601 & 0.6413 \\\\ \\hline zD4 \t\t\t& 6.09 -- 6.73 & 0.24 & 150.2 & 0 & 0 & 0.1869 \\\\ & 7.25 -- 8.00 & 0.33 & 100.8 & 0 & 0.0689 & 0.5326 \\\\ \\hline P34.z.4809\t& 3.60 -- 7.40 (Optical) & 0.70 & 19.45 & 0.8268 & 0.9188 & 0.9856  \\\\ & 7.42 - 19.40 (NIR) & 0.30 & 35.21 & 0.8191 & 1.2917 & 1.7287  \\\\ \\hline ERS.YD2\t\t& 7.42 -- 19.40 (NIR)  & 0.86 & 67.19 & 0.6529 & 1.1551 & 1.4293 \\\\ \\hline HUDF.YD3\t& 7.42 -- 19.40  (NIR) & 0.98 & 241.6 & 0.0550 & 0.1881 & 0.6516 \\\\ \\hline UDF092y-03751196d & 7.42 -- 19.40 (NIR) & 0.996 &  28.64 & 0.8400 & 0.9374 & 0.9872 \\\\ \\end{tabular} \\label{table:fractions} \\end{table*}   \\begin{table*} \\caption{This table shows the total effective number of sampled galaxies with an EW upper limit lower than a set threshold.  We present these figures separately for $z$-drops, $Y$-drops and both $z$-drops and $Y$-drops combined.  The figures in this table are for an unresolved line.} \\begin{tabular}{| c | c | c | c | c |} \\hline & Average Redshift & & $N_{\\mathrm{eff}}=\\sum$ Frac$_{z} \\times$ Frac$_{\\mathrm{EW < thres}}$  & \\\\ & & & & \\\\ & & EW$_{\\mathrm{thres}}=50$\\,\\AA & EW$_{\\mathrm{thres}}=75$\\,\\AA & EW$_{\\mathrm{thres}}=120$\\,\\AA \\\\ \\hline z-drops & 7.0 & 1.242 & 1.903 & 2.963 \\\\ \\hline $Y$-drops & 8.5 & 1.452 & 2.111 & 2.851 \\\\ \\hline $z$-drops \\& $Y$-drops & 7.8 & 2.694 &  4.014 & 5.814 \\\\ \\hline \\end{tabular} \\label{table2} \\end{table*}     \\begin{table*} \\caption{This table for all objects targeted by our spectroscopy shows the redshift range spanned by our data for Lyman-$\\alpha$ in Column 2.  Column 3 shows the fractional probability that a galaxy drawn from the dropout sample would fall within the spectral coverage of that particular spectrograph setup.  Column 4 gives the median EW for each object for the most probable redshift range (see Figure \\ref{fig:sims}).  The remaining three columns tabulate the fraction of our spectroscopy that has EW limits lower than our chosen threshold (50\\,\\AA, 75\\,\\AA\\, and 120\\,\\AA\\, respectively.) The figures in this table are for a 200km/s line.} \\begin{tabular}{ | c | c | c | c | c | c | c | } \\hline Object & z-Range spanned & Frac$_{z}$ & Median EW & Frac$_{\\mathrm{EW < 50\\AA}}$ &  Frac$_{\\mathrm{EW < 75\\AA}}$ & Frac$_{\\mathrm{EW < 120\\AA}}$\\\\ % & by data (for Ly-$\\alpha$)& & & & & \\\\ \\hline \\hline zD1 \t\t\t& 5.82 -- 6.40 & 0.01 & 107.4 & 0 & 0 & 0 \\\\ & 6.08 -- 6.74 & 0.24 & 476.2 & 0 & 0 & 0 \\\\ & 6.79 -- 7.45 & 0.50 & 69.38 & 0.1958 & 0.5460 & 0.8672 \\\\ & 7.25 -- 8.05 & 0.36 & 167.3 & 0 & 0 & 0.1016 \\\\ \\hline zD2\t\t\t& 6.09 -- 6.73 & 0.24 & 222.5 & 0 & 0 & 0 \\\\ & 7.25 -- 8.00 & 0.33 & 149.3 & 0 & 0 & 0.2177 \\\\ \\hline zD3\t\t\t& 6.09 -- 6.73 & 0.24 & 235.9 & 0 & 0 & 0 \\\\ & 7.25 -- 8.00 & 0.33 & 158.4 & 0 & 01 & 0.1062 \\\\ \\hline zD4 \t\t\t& 6.09 -- 6.73 & 0.24 & 294.4 & 0 & 0 & 0 \\\\ & 7.25 -- 8.00 & 0.33 & 197.7 & 0 & 0 & 0 \\\\ \\hline P34.z.4809\t& 3.60 -- 7.40 (Optical) & 0.70 & 38.87 & 0.5759 & 0.6811 & 0.8617  \\\\ & 7.42 -- 19.40 (NIR) & 0.30 & 58.65 & 0.2650 & 0.7361 & 1.2627  \\\\ \\hline ERS.YD2\t\t& 7.42 -- 19.40 (NIR)  & 0.86 & 112.4 & 0.1583 & 0.4344 & 1.0600 \\\\ \\hline HUDF.YD3\t& 7.42 -- 19.40  (NIR) & 0.98 & 400.3 & 0 & 0 & 0.0431 \\\\ \\hline UDF092y-03751196d & 7.42 -- 19.40  (NIR) & 0.996 &  50.63 & 0.5331 & 0.7838 & 0.9181 \\\\ \\end{tabular} \\label{table:fractions_200kms} \\end{table*}  \\begin{table*} \\caption{This table shows the total effective number of sampled galaxies with an EW upper limit lower than a set threshold.  We present these figures separately for $z$-drops, $Y$-drops and both $z$-drops and $Y$-drops combined. The figures in this table are for a 200km/s line.} \\begin{tabular}{| c | c | c | c | c |} \\hline & Average Redshift & & $N_{\\mathrm{eff}}=\\sum$ Frac$_{z} \\times$ Frac$_{\\mathrm{EW < thres}}$  & \\\\ & & & & \\\\ & & EW$_{\\mathrm{thres}}=50$\\,\\AA & EW$_{\\mathrm{thres}}=75$\\,\\AA & EW$_{\\mathrm{thres}}=120$\\,\\AA \\\\ \\hline z-drops & 7.0 & 0.581 & 0.971 & 1.559 \\\\ \\hline $Y$-drops & 8.5 & 0.667 & 1.154 & 1.868 \\\\ \\hline $z$-drops \\& $Y$-drops & 7.8 & 1.248 & 2.125 & 3.427 \\\\ \\hline \\end{tabular} \\label{table2_200kms} \\end{table*}"
    },
    "1208/1208.5455_arXiv.txt": {
        "abstract": "Cyg X-3 is a highly interesting accreting X-ray binary, emitting from the radio to high-energy gamma-rays. It consists of a compact object wind-fed by a Wolf-Rayet (WR) star, but the masses of the components and the mass-loss rate have been a subject of controversies. Here, we determine its masses, inclination, and the mass-loss rate using our derived relationship between the mass-loss rate and the mass for WR stars of the WN type, published infrared and X-ray data, and a relation between the mass-loss rate and the binary period derivative (observed to be $>0$ in Cyg X-3). Our obtained mass-loss rate is almost identical to that from two independent estimates and consistent with other ones, which strongly supports the validity of this solution. The found WR and compact object masses are $10.3_{-2.8}^{+3.9} \\msun$, $2.4_{-1.1}^{+2.1} \\msun$, respectively. Thus, our solution still allows for the presence of either a neutron star or a black hole, but the latter only with a low mass. However, the radio, infrared and X-ray properties of the system suggest that the compact object is a black hole. Such a low-mass black-hole could be formed via accretion-induced collapse or directly from a supernova. ",
        "introduction": "\\label{intro}  Cyg X-3 is an X-ray binary possessing a number of unique and highly interesting characteristics. It is the brightest radio source among X-ray binaries \\citep{mccollough99}, showing extremely strong radio outbursts and resolved jets (e.g., \\citealt*{marti01,m01}). It is, so far, the only X-ray binary that is certainly powered by accretion for which emission of high-energy \\g-rays has been unambiguously confirmed \\citep{agile,fermi09}. It is also the only known X-ray binary in the Galaxy with a Wolf-Rayet (WR) donor \\citep*{v92,v96,v93,fender99}. Furthermore, its orbital period of $P\\simeq 0.2$ d is unusually short for a high-mass binary, which indicates a past spiral-in episode during a common-envelope evolutionary stage. Given the value of $P$ and estimates of the masses and mass-loss rate, the compact object orbits within the WR photosphere.  In spite of the discovery of Cyg X-3 already in 1966 \\citep{giacconi67}, the nature of its compact object has remained unknown; it may be either a neutron star (NS) or a black hole (BH, e.g., \\citealt*{hanson00,vilhu09}, hereafter V09). This issue is of great interest since Cyg X-3 is a likely progenitor of a double compact system, after its WR star explodes as a supernova. If it contains an NS, it may be a progenitor of double NS systems like PSR B1913+16 \\citep{ht75}. That pulsar has a relatively short spin period, 59 ms, and a weak magnetic field, $\\simeq 2\\times 10^{10}$ G. It might have been recycled by accretion in a system like Cyg X-3, which both weakened the initially strong magnetic field and spun it up before the donor exploded and formed a second NS (see, e.g., \\citealt{vdh95} and references therein). On the other hand, if it contains a BH, it may lead to formation of a BH-NS or double BH system, which merger can then be detectable in gravitational waves \\citep{paper2}. Evolutionary calculations and population synthesis for Cyg X-3 were performed by \\citet{lommen05}. They found that either a BH or an NS could be present, although they favoured the BH.  Here, we self-consistently determine the binary parameters and the mass-loss rate, $\\dot M$, of the donor. We use published results for the radial velocities of the two components, and a relationship between the rate of the period increase, $\\dot P$, $\\dot M$, and the total mass of the system. We close the equations by finding a tight correlation between $\\dot M$ and the WR mass, $M_{\\rm WR}$ (for the WN stellar type, identified to be present in Cyg X-3). ",
        "conclusions": ""
    },
    "1208/1208.4855_arXiv.txt": {
        "abstract": "{We present a new prescription for analysing cosmological perturbations in a more-general class of scalar-field dark-energy models where the energy-momentum tensor has an imperfect-fluid form. This class includes Brans-Dicke models, $f(R)$ gravity, theories with kinetic gravity braiding and generalised galileons. We employ the intuitive language of fluids, allowing us to explicitly maintain a dependence on physical and potentially measurable properties. We demonstrate that hydrodynamics is not always a valid description for describing cosmological perturbations in \\emph{general }scalar-field theories and present a consistent alternative that nonetheless utilises the fluid language. \\\\ We apply this approach explicitly to a worked example: k-\\emph{essence }non-minimally coupled to gravity. This is the simplest case which captures the essential new features of these imperfect-fluid models. We demonstrate the generic existence of a new scale separating regimes where the fluid is perfect and imperfect. We obtain the equations for the evolution of dark-energy density perturbations in both these regimes. The model also features two other known scales: the Compton scale related to the breaking of shift symmetry and the Jeans scale which we show is determined by the speed of propagation of small scalar-field perturbations, i.e. causality, as opposed to the frequently used definition of the ratio of the pressure and energy-density perturbations.} ",
        "introduction": "}  Many models, better or worse motivated by fundamental physics, have been proposed as an alternative to the cosmological constant for providing a mechanism to accelerate the expansion of the universe at late times. Perhaps the simplest, and therefore most studied, are models which feature a single extra scalar degree of freedom, which is treated as classical. The archetypal example of this dark energy (``DE'') are quintessence models \\cite{Ratra:1987rm,Wetterich:1987fm}, where a canonical scalar field rolls in a potential resulting a time-varying equation of state and allowing for (small) inhomogeneities in the dark-energy fluid. On the other hand, $f(R)$ gravity models \\cite{Carroll:2003wy}, where the Einstein-Hilbert action for gravity is modified, allow for a different response of the metric to the presence of matter from that in general relativity (``GR'') and usually are discussed as examples of modified gravity. However, these models can also be reformulated as a particular subclass of Brans-Dicke theories featuring a single extra scalar, albeit non-minimally coupled to gravity \\cite{Brans:1961sx,Chiba:2003ir}. It was quickly understood that whereas the background expansion history can be indistinguishable between various classes of these models, the evolution of dark-matter perturbations can differ significantly and therefore formation of large-scale structure provides the key to understanding the role played by these theories, if any, in accelerating the universe \\cite{Ishak:2005zs}.  k-\\emph{essence} was introduced as an extension of quintessence models to non-canonical kinetic terms \\cite{ArmendarizPicon:1999rj,ArmendarizPicon:2000ah,ArmendarizPicon:2000dh} and can be interpreted as a perfect fluid with a speed of sound different from that of light \\cite{Garriga:1999vw}. This allows for a Jeans scale smaller than the cosmological horizon and therefore for DE which can cluster significantly. In particular, the discovery of generic attractors where the equation of state approaches that of the vacuum while the sound speed tends to zero \\cite{ArkaniHamed:2003uy} also spurred the discussion of such models as describing both dark energy and dark matter through a single degree of freedom \\cite{Scherrer:2004au,Bertacca:2007ux,Lim:2010yk}. One tends to think of these dark-energy models as perfect fluids, obeying hydrodynamics and studies the evolution of linear perturbations by considering the conservation of the perturbed energy-momentum tensor \\cite{Bardeen:1980kt,Ma:1995ey,Hu:1998tj}.  On the other hand, when studying large-scale structure in $f(R)$ theories \\cite{Song:2006ej,Tsujikawa:2007xu}, one tends to think in terms of a modification of the response of the two scalar gravitational potentials, $\\Phi$ and $\\Psi$, to matter perturbations. In particular the generic result for this class of theories is that the lensing potential, $\\Phi-\\Psi$, is ``unaffected'' and depends only on the matter perturbations and an evolving Newton's constant, while the modification of gravity creates anisotropic stress giving $2\\Phi+\\Psi=0$.% \\footnote{For an investigation of the relation between anisotropic stress and cosmological evolution in non-linear gravity models see ref \\cite{Saltas:2010tt}.% } These relationships are simple and therefore are frequently generalised in order to constrain such theories from data. A large number of phenomenological parameterisations of models of modified gravity exist \\cite{Linder:2005in,Knox:2005rg,Ishak:2005zs,Kunz:2006ca,Laszlo:2007td,Caldwell:2007cw,Amendola:2007rr,Pogosian:2010tj} and usually focus on directly modifying the relationships between the potentials and matter perturbations without necessarily asking what this implies for the underlying model. Considering the lack of well-motivated models, this might not be considered a disadvantage. A different approach, based on symmetries of the Einstein equations without a prior reference to a particular dark-energy model, was employed in ref.~\\cite{Battye:2012eu}, while a parameterised post-Friedmann approach was proposed in \\cite{Hu:2007pj} and further developed in \\cite{Baker:2011jy}.  $f(R)$ models and this sort of modifications of gravity are quite constrained by Solar-System tests as a result of their being equivalent to fifth forces. One can find viable models \\cite{Hu:2007nk,Starobinsky:2007hu,Amendola:2006we} which can evade the constraints through the chameleon effect \\cite{Khoury:2003aq,Khoury:2003rn}, or one can propose that the extra scalar does not couple to baryons, but only to dark matter \\cite{Wetterich:1994bg,Amendola:1999er}. These types of interacting DE models can have dynamics on large scales very similar to modified gravity models, since the contribution of baryons to the energy density is very subdominant. However, philosophically they are no longer ``modified gravity'' since the coupling is not universal --- it no longer couples purely to the energy-momentum tensor. On the other hand, the full breadth of such models is much wider than for $f(R)$ theories, which are but a subclass. When studying interacting dark energy, one tends to treat the scalar field differently from matter: the equation of motion for the scalar is usually parameterised and solved, while the dark matter and baryons are treated as a hydrodynamical fluid.  Another class of scalar-field models was discovered by considering gravity modifications arising from a potential higher-dimensional embedding of our universe \\cite{Dvali:2000hr,Deffayet:2000uy,Gabadadze:2009ja}. In the decoupling limit \\cite{Luty:2003vm,Nicolis:2004qq,Gabadadze:2006tf}, this DGP gravity reduces to a scalar theory, where the scalar Lagrangian contains second derivatives of the scalar field and therefore is not in the k-\\emph{essence} class. A further generalisation of such scalar-field theories featuring second derivatives was then obtained in refs \\cite{Nicolis:2008in,Nicolis:2009qm} and named the \\emph{galileon}. Such theories are not consistent on general curved backgrounds, and the appropriate covariantisation was obtained in \\cite{Deffayet:2009wt,Deffayet:2009mn}. Kinetic gravity braiding, a generalisation of the DGP decoupling limit, was independently found in refs \\cite{Deffayet:2010qz,Kobayashi:2010cm} and soon it was proven that the most general theory of a single scalar with no more than second derivatives is a similar generalisation of the galileons \\cite{Deffayet:2011gz}. Indeed these generalised galileons were discovered a long time ago \\cite{Horndeski:1974} and largely ignored. The higher-dimensional origin of the galileon theories was studied in refs \\cite{deRham:2010eu,Hinterbichler:2010xn,Goon:2011qf,Goon:2011uw}, while the galileon theories also have appeared in the decoupling limit of consistent massive-gravity as actions for the helicity-0 mode of the massive graviton \\cite{deRham:2010ik,deRham:2010kj,Hassan:2011vm}.  These theories all contain second derivatives in the action and therefore have energy-momentum tensors (``EMT'') which do \\emph{not }have the form of a perfect fluid \\cite{Pujolas:2011he,Gao:2011mz} and therefore provide for new phenomenology. For example, the null-energy condition (``NEC'') can be violated stably, resulting in explicitly ghost-free phantom DE models or even permitting the construction of initial phases of evolution of the universe alternative to inflation \\cite{Creminelli:2010ba,Qiu:2011cy,Easson:2011zy}. Equations for linear cosmological perturbations for the most general of such theories were derived in \\cite{Gao:2011qe,DeFelice:2011hq}.\\textbf{ }Indeed, models with stabilised ``phantom-crossing'' and subsequent evolution were discussed from the point of view of the effective action for perturbations already in \\cite{Creminelli:2006xe,Creminelli:2008wc}: adding higher-derivative operators acts to stabilise configurations violating the NEC.\\\\   What strikes in the above discussion is the multitude of approaches that are used to study perturbations, each chosen depending on the underlying Lagrangian for the scalar. However, all the models discussed here are structurally extremely similar. A common framework for describing perturbations would allow for a single consistent method of constraining the class of models of DE comprising a single scalar. One could argue that solving the equation of motion for the scalar provides exactly such a framework. However, if we assume that the DE cannot be studied directly, but only through its effect on the gravitational field in which baryons and light propagate, the equation of motion contains too much information. Gravity couples only to the EMT and therefore will mostly depend on the fluid properties of the DE and not on the precise value of the scalar field.  The case of k-\\emph{essence} is a pertinent example. Any particular choice within this class of Lagrangians gives a particular expansion history for the cosmological background, with a particular evolution of the equation-of-state parameter, $w$. This solution (based on the choice of Lagrangian) then also determines the speed of propagation of the scalar-field perturbations, $c_{\\text{s}}^{2}$. Given the history of these two variables, the evolution of linear energy-density perturbations of a k\\emph{-essence }DE\\emph{ }is \\emph{completely} determined at all scales. Any k-\\emph{essence} model which happens to have the same history of evolution of $w$ and $c_{\\text{s}}^{2}$ will have exactly the same evolution of linear perturbations, even if the two Lagrangians are very different and only happen to have coincided in this way on two individual solutions. There is therefore no way to reconstruct a Lagrangian even limited to this single class of theories, even if we ignore the fact that the data that can be obtained will only ever be limited to some maximal precision and we will only be able to measure the properties of the DE indirectly. Unless DE interacts directly with baryonic matter in a way that is not equivalent to redefining a gravitational metric, we will only be able to measure the properties of the DE perturbations indirectly by looking at scale-dependent modifications of the DM power spectrum. Given this, the best we can hope for is to measure $w$ through its effect on geometry and the sound speed of DE by detecting some sort of scale of transition in behaviour that would be consistent with a Jeans length \\cite{Sapone:2009mb,Sapone:2010uy}.  A similar situation occurs when dark energy is described by the class of $f(R)$ theories. Here, there appears to be only one scale, that of the (time-varying) Compton wavelength of the scalar, the chameleon effect allowing for the restoration of GR in dense regions through non-linearities notwithstanding. On both sides of this scale, the evolution of linear perturbations is completely determined by the fact that we have chosen to consider this limited class of Lagrangians, just as it was for k-\\emph{essence}. Indeed, specifying the evolution of $w$ and of the Compton scale gives us the totality of the information that can be gleaned from the background and the linear perturbations.  In fact, the standard discussion of $f(R)$ gravity neglects two more scales which are implicitly contained in the theory. Firstly, there is still the Jeans scale. However, in this case, the speed of sound is equal to that of light and therefore the Jeans scale lies at the cosmological horizon and is thus unobservable. Secondly, there exists a new scale controlling whether the fluid is perfect of imperfect, which will be the subject of much of this paper. $f(R)$ models are a limit where this scale lies at infinity and therefore it is also unobservable there. \\\\   The main aim of this paper is to provide a prescription for describing linear perturbations which as far as possible depends only on physical properties of the model, such as the equation of state or the sound speed, which one could hope to measure. Our approach is to describe the DE through the conservation equations for its EMT, rather than the equation of motion. This provides the evolution equations for the energy-density perturbations which are then directly connected to the gravitational potentials. For any theory, these equations can be solved provided we supply two closure relations: relationships of the pressure perturbation and anisotropic stress to the energy-density perturbation. The solutions for linear perturbations of any fluid are always fully determined by these two functions.% \\footnote{Some EMT non-conservation issues related to the evolution of the effective Planck mass notwithstanding.% } The question we answer is how to properly obtain these relationships for a model featuring a single scalar degree of freedom.  We show explicitly that even for general k-\\emph{essence} models these closure relations, which are well known \\cite{Kunz:2006wc,Christopherson:2008ry}, are not hydrodynamical and one should never a priori assume hydrodynamical behaviour in more general cases. However, on any time slice all components of the EMT are determined by the values of the degrees of freedom at that moment in time. Thus, for any class of models based on degrees of freedom with classical equations of motion, we can obtain such closure relations. Indeed these relations define the totality of the properties of the dark energy at the linear level of perturbations. We provide a form of these closure relations for any scalar-field theory and we explicitly calculate them for k-\\emph{essence} coupled non-minimally to gravity. This particular class contains both k-\\emph{essence }and $f(R)$ theories and as a result of the non-minimal coupling to gravity features second-derivatives of the scalar in the EMT in the Jordan frame. In this sense, it mimics the form of the EMT possessed by more-general galileon theories and therefore is a useful example.  We show that in general scalar-field models we have three scales: the Jeans scale, the Compton scale and a new scale determining whether the fluid is in a perfect or imperfect regime. This is \\emph{in addition }to the cosmological-horizon scale. These three scales are essentially independent and can lie in any order and are determined by the Lagrangian class and the particular background solution. This new transition scale appears in all scalar-field models more general than k-\\emph{essence} and is determined by the relative importance of the second-derivative terms and the k-\\emph{essence} terms in the perturbed EMT. For example, it is only inside this scale that the DE fluid can carry anisotropic stress at all. This scale is not visible in the equation of motion for the scalar, but only appears when the EMT is considered.  As an example of the power of our prescription, we obtain the equations for the evolution of energy-density perturbations for the non-minimally coupled \\emph{k-essence }DE model and solve them analytically under the assumption that the background evolution exhibits scaling behaviour (i.e. has a constant equation-of-state parameter).\\\\   The paper is structured as follows: we start in section \\ref{sec:FluidFormalism} by describing a formalism for decomposing the EMT into the fluid variables covariantly. This description allows us to obtain exact forms for, among others, the energy density and pressure in terms of the scalar-field, which can then be perturbed directly. In this way, we obtain the conservation equations for the linear perturbations of the EMT in section \\ref{sub:EvolEqLin}. In section \\ref{sub:hydro} we discuss what hydrodynamics implies for the relationship between pressure and energy density in perfect fluids, and by deriving this relationship for general k-\\emph{essence }models in section \\ref{sub:k-ess} we show that the hydrodynamical relations are not obeyed. We then propose a form for these closure relations which would be obeyed by any DE model comprising a single scalar field in section \\ref{sub:GenPerts}. In section \\ref{sec:BDK}, we turn to a non-minimally coupled k\\emph{-essence }as worked\\emph{ }example. We discuss the general properties of this theory, before turning to calculate the closure relations between the pressure, energy density and anisotropy perturbations in section \\ref{sub:DEonly_perts}. We demonstrate explicitly how to calculate them and show that they are of the form contained within the general parameterisation introduced earlier. We explicitly demonstrate the existence of the three scales in the problem and derive the equation for the evolution of the density contrast in these models. We summarise our findings in section \\ref{sec:Conc}. ",
        "conclusions": "}  In this paper we have proposed a unified prescription for studying the subhorizon evolution of linear perturbations in theories of dark energy comprising a single scalar, including, for example, modified gravity models such as $f(R)$ and other models containing second-derivative terms in their EMTs. These EMTs are in general not of perfect-fluid form.  k-\\emph{essence} is the most general scalar theory that has an EMT of perfect-fluid form. The evolution of linear cosmological perturbations in this class of models is already well understood. However, any further modifications to the scalar's Lagrangian, such as galileon terms or even a simple non-minimal coupling to gravity, generate terms including second derivatives of the scalar in the EMT. The presence of the second-order derivatives makes it impossible to express the EMT in perfect-fluid form. A standard method for dealing with such models, as carried out in the literature, is to turn to the equation of motion for the scalar, usually concentrating on the limit of negligible time derivatives for the scalar field. We have argued that this is in fact on one hand too specific. For example, for k\\emph{-essence} any two background solutions with the same history of equation-of-state parameter $w$ and sound speed $c_{\\text{s}}^{2}$ will give exactly the same evolution of linear perturbations, irrespective of the actual Lagrangian, as eq.~(\\ref{eq:d-k-subdom-1}) shows. One may of course prefer one of the Lagrangians over another on the basis of naturalness, but this is not an observable if we only have access to the cosmological background and the effect of linear density perturbations on gravitational potentials. On the other hand: the neglecting of the time derivatives could prove dangerous, if the coefficient of the friction term in the density-contrast evolution equation turns out to be small enough to allow for modes which grow sufficiently quickly on a particular cosmological background. We have, for example, explicitly demonstrated that in the non-minimally coupled k-\\emph{essence} model, the DE energy-density contrast during the DE-domination era evolves in such a way so as to cause the gravitational potential to grow, eq.~(\\ref{eq:impfmodes}). These modes should dominate the gravitational potential at late-enough times on at least some of the scales.  Our prescription allows for the study of such general models in terms of the intuitive fluid language instead of dealing directly with the equation of motion for the scalar field. Since the fluid variables such as pressure and energy density are just appropriately reorganised components of the EMT, they will contain all the information to which the gravitational field is sensitive.  As is well known, one can obtain the evolution equation for the energy-density perturbations from the perturbed conservation equations for the EMT. However, two closure relations need to be provided to do this: we need to relate both the pressure perturbation and the anisotropic-stress perturbation to the energy-density perturbation. The main result of this paper is the method for obtaining these closure relations, based on considering the actual degrees of freedom on a spatial hypersurface.  As we have explicitly shown in our example, this method produces closure relations that are not those of perfect-fluid hydrodynamics. One \\emph{cannot }assume that the relationship between the pressure and energy-density perturbations is hydrodynamical, $\\delta\\mathcal{P}=c_{\\text{H}}^{2}\\delta\\mathcal{E}$, with $c_{\\text{H}}$ a hydrodynamical sound speed for pressure waves, as is frequently done. We have explained in detail in sections \\ref{sub:hydro} and \\ref{sub:k-ess} why even general k-\\emph{essence} models should not be thought of as a hydrodynamical fluid. The result is that the closure relations for general scalar-field theories take the somewhat complicated form (\\ref{eq:ClosureParam}) and cannot be reduced to a simple relationship.  One may argue that this hydrodynamical relation between $\\delta\\mathcal{P}$ and $\\delta\\mathcal{E}$ is the definition of the sound speed for pressure waves, $c_{\\text{H}}$. However, in general this is a not a physically meaningful quantity: in the example we have studied, where the anisotropic stress does not vanish, the Jeans length is not at all related to the quantity $c_{\\text{H}}$, but is determined by the speed of propagation of perturbations of the scalar field, $c_{\\text{s}}$, and therefore causality. As we explicitly show, $c_{\\text{s}}$ is obtained by considering the effective wave operator in the scalar perturbed equation of motion and is not immediately apparent in the expressions for the pressure perturbation. Even when $c_{\\text{s}}$ is simple and constant, the complexity of the closure relations we have derived shows that $c_{\\text{H}}$ could be an arbitrary function of time and space and therefore not at all a useful parameter for the description of the perturbations.\\\\   In order to aid the calculation of the closure relations, we have introduced a book-keeping parameterisation for them, suitable for a wide range of models containing a single scalar degree of freedom. This has allowed us to derive a general equation for the evolution of DE density perturbations in terms of these parameters which can be calculated for any particular class of scalar Lagrangians. For the purpose of illustration, we have shown explicitly how to calculate the coefficients of the closure relations (\\ref{eq:ClosureParam}) in a toy class of models, presenting the results in table \\ref{tab:ClosureParams}. All the parameters which are non-vanishing for this class of Lagrangians are expressible in terms of physical quantities: $w$, $c_{\\text{s}}^{2}$, the coupling to external matter $\\beta$ and the Compton term $M_{C}^{2}$, each of which is at most a function of time but never scale. Non-minimally coupled k-\\emph{essence} was chosen as the toy model since it is the simplest one to contain the majority of the features which are present in more-general scalar theories such as galileons. This prescription can be directly applied in the study of those more-complex models.  All models containing terms with second derivatives in the EMT also contain a new scale $k_{\\text{T}}$, eq.~(\\ref{eq:kT}), determined by the relative contribution of the k-\\emph{essence }and second-derivative terms in the perturbations. In our toy example, as we move across this scale, the DE fluid transitions from k-\\emph{essence}-like perfect-fluid behaviour at large subhorizon scales to an imperfect fluid carrying anisotropic stress at small scales.\\textbf{ }Interestingly, the scale $k_{\\text{T}}$ does \\emph{not} appear in the perturbed equation of motion for the scalar itself, just in the EMT. On both sides of the transition scale, the scalar-field perturbation $\\delta\\phi$ evolves in the same manner, at least provided there isn't a scale dependence in any sources on the external-matter side. The changes in the pressure and anisotropy closure relations as we cross the scale $k_{\\text{T}}$ reflect simply that given the same evolution in time of $\\delta\\phi$, the density perturbation $\\delta\\mathcal{E}$ must evolve differently, since the \\emph{k-essence }and the second-derivative terms have a completely different functional dependence on $\\delta\\phi$. Since it is $\\delta\\mathcal{E}$ that sources the gravitational potential, this transition scale is physical. In addition to the transition scale, the dark energy will feature the known Jeans scale, controlled by the speed of propagation of small perturbations of the scalar field and a Compton scale related to the mass of the field. We have not studied superhorizon behaviour and the transition towards the horizon, but since the new phenomenology is related to the presence of second derivatives in the EMT, it will be irrelevant at large enough scales, where the standard k-\\emph{essence }behaviour is restored. Indeed sufficiently superhorizon the curvature perturbation is conversed in the standard way \\cite{Gao:2011mz}.  What our discussion shows is that it is more appropriate to think of the anisotropic stress as being related through the closure relation (\\ref{eq:ClosureParam}) to the density perturbation of the DE fluid rather than directly to the gravitational potentials. In general, the gravitational potential will also have contributions from other fluids present in the universe. The density perturbations in any particular fluid evolve according to the conservation equations for that fluid, and therefore in principle are independent dynamical variables. In our example model, in the imperfect regime, the closure relations imply that $k^{2}\\delta\\pi/a^{2}\\simeq\\delta\\mathcal{E}$, \\emph{irrespective} of the configuration of the external matter. It is this property of the fluid that makes the lensing potential independent of the perturbation in the scalar field, eq.~(\\ref{eq:lensing}). And it is only this relation that is in general independent of the other constituents of the universe.  The only reason why in $f(R)$ models this scale $k_{\\text{T}}$ is not seen, is that they are a special subclass of the non-minimally coupled k\\emph{-essence }models\\emph{ }where the k-\\emph{essence} terms are absent. Thus the DE fluid is imperfect at all scales. The Compton scale determines whether for a particular mode the energy-density perturbations are large or small. At scales inside the Compton radius they are large and therefore the anisotropic stress is large. Outside --- they are small and so is the anisotropic stress.  We should stress that every \\emph{class }of Lagrangian terms which generates second derivatives in the EMT (e.g. kinetic gravity braiding, etc.) will produce such a transition scale determined by the relative size of its coefficient in the background solution. Each class of terms will presumably cause slightly different physical effects and therefore each of these transitions should be potentially observable.  Finally, we should reiterate that in order to simplify the exposition, we have mostly neglected in this presentation the effect of external matter. This has allowed us to concentrate on understanding the properties of the DE fluid itself. As a result of the non-minimal coupling to gravity, the DE and dark matter are coupled in this model. As we have previewed in the closure relations (\\ref{eq:ClosureParam}), the coupling to dark matter causes the DM perturbation to appear in the expression for $\\delta\\mathcal{P}$ and therefore in the Jeans term. The DM perturbations will in this way provide a scale-dependent source in the evolution equation for the energy-density perturbation $\\delta\\mathcal{E}$, dominant in the imperfect regime. This ensures that the density perturbations for DE are large in the imperfect regime, providing a significant modification to the gravitational potentials in which dark matter propagates.  The presence of an external energy density provides an additional complication: the way we defined our DE EMT means that in this case it is not conserved whenever the effective Planck mass evolves, i.e. in models with non-minimal coupling to gravity. As we show in the expressions for the conservation of the EMT (\\ref{eq:EnConsPara}) and (\\ref{eq:MomConsPara}), there are non-conservation terms which depend on the scalar and the dark-matter configuration. Therefore, we will have to provide a similar closure relation to (\\ref{eq:ClosureParam}), one for each of energy and momentum conservation equations. We will return to this also feature-rich discussion in a separate work."
    },
    "1208/1208.6582_arXiv.txt": {
        "abstract": "{Simple models fail to describe the observed spectra of X-ray dim isolated neutron stars (XDINSs). Interpretation of these spectra requires detailed studies of radiative properties in the outermost layers of neutron stars with strong magnetic fields. Previous studies have shown that the strongly magnetized plasma in the outer envelopes of a neutron star may exhibit a phase transition to a condensed form. In this case thermal radiation can emerge directly from the metallic surface without going through a gaseous atmosphere or, alternatively, it may pass through a ``thin'' atmosphere above the surface. The multitude of theoretical possibilities complicates modeling the spectra and makes it desirable to have analytic formulae for constructing samples of models without going through computationally expensive, detailed calculations. }{ The goal of this work is to develop a simple analytic description of the emission properties (spectrum and polarization) of the condensed, strongly magnetized surface of neutron stars. }{ We have improved the method of van Adelsberg et al.~(2005) for calculating the spectral properties of condensed magnetized surfaces. Using the improved method, we calculate the reflectivity of an iron surface at magnetic field strengths $B\\sim10^{12}$~G -- $10^{14}$~G, with various inclinations of the magnetic field lines and radiation beam with respect to the surface and each other. We construct analytic expressions for the emissivity of this surface as functions of the photon energy, magnetic field strength, and the three angles that determine the geometry of the local problem. Using these expressions, we calculate X-ray spectra for neutron stars with condensed iron surfaces covered by thin partially ionized hydrogen atmospheres. }{ We develop simple analytic descriptions of the intensity and polarization of radiation emitted or reflected by condensed iron surfaces of neutron stars with strong magnetic fields typical for isolated neutron stars. This description provides boundary conditions at the bottom of a thin atmosphere, which are more accurate than previously used approximations. The spectra calculated with this improvement show absorption features different from those in simplified models. }{ The approach developed in this paper yields results that can facilitate modeling and interpretation of the X-ray spectra of isolated, strongly magnetized, thermally emitting neutron stars. } ",
        "introduction": "\\label{sect:intro}  Recent observations of neutron stars have provided a wealth of valuable information, but they have also raised many new questions. Particularly intriguing is the class of radio-quiet neutron stars with thermal-like spectra, commonly known as X-ray dim isolated neutron stars (XDINSs), or the Magnificent Seven (see, e.g., reviews by \\citealt{Haberl07} and \\citealt{Turolla09}, and references therein). Some of them (e.g., RX J1856.5$-$3754) have featureless spectra, whereas others (e.g., RX J1308.6+2127 and RX J0720.4$-$3125) have broad absorption features with energies $\\sim0.2$\\,--\\,2 keV. In recent years, an accumulation of observational evidence has suggested that XDINSs may have magnetic fields $B\\sim10^{13}$\\,--\\,$10^{14}$~G and be related to magnetars (e.g., \\citealt{Mereghetti08}).  For interpretation of the XDINS spectra, it may be necessary to take the phenomenon of ``magnetic condensation'' into account. The strong magnetic field squeezes the electron clouds around the nuclei and thereby increases the binding and cohesive energies (e.g., \\citealp{MedinLai06}, and references therein). Therefore XDINSs may be ``naked,'' with no appreciable atmosphere above a condensed surface, as first conjectured by \\citet{ZaneTD02}, or they may have  a relatively thin atmosphere, with the spectrum of outgoing radiation affected by the properties of the condensed surface beneath the atmosphere, as suggested by \\citet{MotchZH03}.  Reflectivities of condensed metallic surfaces in strong magnetic fields have been studied in several papers \\citep{Brinkmann80,TurollaZD04,vanAdelsberg-ea05,PerezAMP05}. \\citet{Brinkmann80} and \\citet{TurollaZD04} neglected  the motion of ions in the condensed matter, whereas \\citet{vanAdelsberg-ea05} (hereafter Paper~I)  and \\citet{PerezAMP05} considered two opposite limiting cases: one that neglects the ion motion (``fixed ions'') and another where the ion response to the electromagnetic wave is treated neglecting the Coulomb interactions between the ions (``free ions''). A large difference between these two limits occurs at photon frequencies below the ion cyclotron frequency, but the two models lead to almost the same results at higher photon energies. We expect that in reality the surface spectrum lies between these two limits (see Paper~I for discussion). The results of Paper~I  and of \\citet{PerezAMP05} are similar to each other, but significantly differ from the earlier results. In particular, \\citet{TurollaZD04} found that collisional damping in the condensed matter leads to a sharp cutoff in the emission at low photon energies, but such a cutoff is absent in Paper~I and \\citet{PerezAMP05}. It is most likely that this difference arises from the ``one-mode'' description for the transmitted radiation adopted by \\citet{TurollaZD04} (see Paper~I for details). All the previous works relied on a complicated method of finding the transmitted radiation modes, originally due to \\citet{Brinkmann80}. We replace it with a more reliable method described below.  \\citet{Ho-ea07} (see also \\citealp{Ho07}) fitted multiwavelength observations of RX J1856.5$-$3754 with a model of a thin, magnetic, partially ionized hydrogen atmosphere on top of a condensed iron surface; they also discussed possible mechanisms of creation of such a thin atmosphere. \\citet{SuleimanovPW09} calculated various models of fully and partially ionized finite atmospheres above a condensed surface including the case of ``sandwich'' atmospheres, composed of hydrogen and helium layers above a condensed surface.  The large variety of theoretical possibilities complicates the modeling and interpretation of the spectra. In order to facilitate this task, \\citet{Suleimanov-ea10} (hereafter Paper~II) suggested an approximate treatment, in which the local spectra together with temperature and magnetic field distributions are fitted by simple analytic functions. Being flexible and fast, this approach is suited to constrain stellar parameters prior to performing more accurate but computationally expensive calculations of model spectra. The reflectivity of the condensed surface was modeled by a simple steplike function, which roughly described the polarization-averaged reflectivity of a magnetized iron surface at $B=10^{13}$~G, but depended neither on the magnetic field strength $B$ nor on the angle $\\varphi$ between the plane of incidence and the plane made by the normal to the surface and the magnetic field lines.  In the present work,  the numerical method of Paper~I and the approximate treatment of Paper~II are refined.  We develop a less complicated and more stable method of calculations and construct more accurate fitting formulae for the reflectivities of a condensed, strongly magnetized iron surface, taking the dependence on arguments $B$ and $\\varphi$ into account. The new fit reproduces the feature near the electron plasma energy, obtained numerically in Paper~I but neglected in Paper~II. Two versions of the fit are presented in Sect.~\\ref{sect:refl} for the models of free and fixed ions discussed in Paper~I. In addition to the fit for the average reflectivity, we present analytic approximations for each of the two polarization modes, which allow us to calculate the polarization of radiation of a naked neutron star. In Sect.~\\ref{sect:atm} we consider the radiative transfer problem in a finite atmosphere above the condensed surface, including the reflection from the inner atmosphere boundary with normal-mode transformations, neglected in the previous studies of thin atmospheres. Conclusions are given in Sect.~\\ref{sect:concl}. In Appendix~\\ref{sect:Improved} we describe the method of calculation for the reflectivity coefficients, which is improved with respect to Paper~I. In Appendix~\\ref{sect:rmj} we describe an analytic model of normal-mode reflectivities at the inner boundary of a thin atmosphere. ",
        "conclusions": "\\label{sect:concl}  We have improved the method of Paper~I for calculating spectral properties of condensed magnetized surfaces. Using the improved method, we have calculated a representative set of reflectivities of a metallic iron surface for the magnetic field strengths $B=10^{12}$~G\\,--\\,$10^{14}$~G. Based on these calculations, we have constructed  analytic expressions for emissivities of the magnetized condensed surface in the two normal modes as functions of five arguments: energy of the emitted X-ray photon $E$, field strength $B$, field inclination $\\theta_B$, and the two angles that determine the photon direction. We have considered the alternative limiting approximations of free and fixed ions for calculating the condensed surface reflectivity.  We have improved the inner boundary conditions for the radiation transfer equation in a thin atmosphere above a condensed surface. The new boundary condition accounts for transformation of normal modes into each other due to reflection from the condensed surface. To implement this condition we suggest a method for calculating reflectivities $R_{MM'}$ in the normal modes used for model atmosphere calculations, based on analytic approximations to the reflectivities.  We computed a few models of thin, partially ionized hydrogen atmospheres to investigate the influence of the new boundary condition on their emergent spectra and temperature structures. The allowance for mode transformations makes the complex absorption feature between $\\Eci$ and $\\EC$ less significant and the atomic absorption feature more important.  Nevertheless, the equivalent widths of this complex absorption feature in the emergent spectra are still significant ($\\approx200$\\,--\\,250 eV) and sufficient to explain the observed absorption features in the spectra of XDINSs. Models of thin atmospheres with inclined magnetic fields are necessary for detailed descriptions of their spectra. We plan to compute such models with vacuum polarization and partial mode conversion in a future paper."
    },
    "1208/1208.3191_arXiv.txt": {
        "abstract": "The Stroke Minimization algorithm developed at the Princeton High Contrast Imaging Laboratory  has proven symmetric dark hole generation using minimal stroke on two deformable mirrors (DM) in series. The windowed approach to Stroke Minimization has proven symmetric dark holes over small bandwidths by using three wavelengths to define the bandwidth of correction in the optimization problem. We address the relationship of amplitude and phase aberrations with wavelength, how this changes with multiple DMs, and the implications for simultaneously correcting both to achieve symmetric dark holes. Operating Stroke Minimization in the windowed configuration requires multiple wavelength estimates. To save on exposures, a single estimate is extrapolated to bounding wavelengths using the established relationship in wavelength to produce multiple estimates of the image plane electric field. Here we demonstrate better performance by improving this extrapolation of the estimate to other wavelengths. The accuracy of the functional relationship will ultimately bound the achievable bandwidth, therefore as a metric these results are also compared to estimating each wavelength separately. In addition to these algorithm improvements, we also discuss a laboratory upgrade and how it can better simulate broadband starlight. We also discuss the possibility of leveraging two DMs in series to directly estimate the electric field over a narrow bandwidth and the challenges associated with it. ",
        "introduction": "\\label{sec:intro} There has been much research into space-based missions for direct imaging of extrasolar terrestrial planets. Two approaches that have been proposed for direct imaging in visible to near-infrared light are occulters, which utilize large external screens to block the starlight while allowing the planet light to pass, and coronagraphs, which use internal masks and stops to change the point spread function of the telescope, creating regions in the image of high contrast where a dim planet can be seen. Observing over a wide band is challenging in both cases and coronagraphs in particular have an extreme sensitivity  to wavefront aberrations generated by the errors in the system optics (occulters are immune to this problem because the starlight never enters the telescope). The chromatic effect of these errors will manifest as aberrations at the image plane that will scale differently in wavelength depending on the type of error and it's location relative to the pupil plane. This necessitates wavefront control algorithms with multi-DM control to correct for the aberrations and allow relaxation of manufacturing tolerances and stability requirements within the observatory to achieve symmetric dark holes in broadband light. In this paper we discuss the challenges associated with wavefront estimation and control, both monochromatic and broadband, in a coronagraphic imager.  We also note that many of these challenges are shared by both space and ground coronagraphs. ",
        "conclusions": "With the new photonic crystal single mode fiber, we have eliminated multimode output of the fiber source. Comparing the most recent results to data collected before the laboratory upgrade, the extrapolation technique was strongly affected by the multimode output at shorter wavelengths. Taking multiple estimates removes some of this sensitivity, but there was still some improvement because the strong wavelength dependence of the input aberrations will limit our ability to achieve simultaneous correction at multiple wavelengths. With the light source upgrade, the two remaining limitations are algorithmic. One is the quality of our DM model.  Since both the estimation and control algorithms incorporate DM surface models and actuator voltage maps, errors in this mapping directly translate into limitations in contrast.  The last source of error is the extent to which the aberrated field over the wavelength band is well represented by the first two terms in a wavelength expansion. Two DMs in series can only correct the $\\lambda$ independent and $1/\\lambda$ terms \\cite{pueyo2007polychromatic}. The theoretical $1\\times10^{-6}$ contrast limit in a $20\\%$ band in the Princeton HCIL is driven by the uncontrollable nominal surface errors of the DMs that will contribute $1/\\lambda^2$ errors due to Fresnel propagation. This points to a stronger requirement on the figure of the DMs to achieve better broadband suppression in symmetric dark holes, and indicates that there would be great value to adding a third DM in series to the optical system since this would make higher order chromatic errors controllable \\cite{pueyo2007polychromatic}."
    },
    "1208/1208.3472_arXiv.txt": {
        "abstract": "We propose that bound, young massive stellar clusters form from dense clouds that have escape speeds greater than the sound speed in photo-ionized gas. In these clumps, radiative feedback in the form of gas ionization is bottled up, enabling star formation to proceed to sufficiently high efficiency so that the resulting star cluster remains bound even after gas removal. We estimate the observable properties of the massive proto-clusters (MPCs) for existing Galactic plane surveys and suggest how they may be sought in recent and upcoming extragalactic observations. These surveys will potentially provide a significant sample of MPC candidates that will allow us to better understand extreme star-formation and massive cluster formation in the Local Universe. ",
        "introduction": "The formation of bound star clusters has become a topic of renewed interest. The Milky Way contains about 150 globular clusters (GCs) with masses from $10^4\\ \\msun$\\ to over $10^6\\ \\msun$\\ and tens of thousands of open clusters containing from 100 to over $10^4$ stars \\citep{PortegiesZwart2010}. While no GCs have formed in the Milky Way within the last 5 Gyr, bound clusters that survive for more than hundreds of crossing times continue to form.  Infrared observations over the last two decades have shown that molecular clouds tend to produce stars in higher surface densities ($\\geq 3~{\\rm stars~pc}^{-2}$) than the field population \\citep{Lada2003}. \\cite{Bressert2010} showed that stars within 500 pc of the Sun form in a smooth continuous distribution and only a minority will dynamically evolve to form bound low-mass stellar clusters ($10^2$ to $10^3 \\msun$). The vast majority of these young clusters are {\\it transient} groups that are bound primarily by the gas in their environment. Thus, while most stars may form in groups, gravitationally bound clusters which remain bound for many crossing-times following dispersal of their natal clump are rare and contain less than 10\\% of the Galactic stellar population. Despite the small number of stars that form in the bound young massive clusters \\citep[YMCs; $\\gtrsim 10^4$ \\msun; ][]{PortegiesZwart2010}, they are important as they shed light on extreme star-formation in the Local Universe and provide insight on how GCs may have formed in the high-redshift universe and in the distant past of the Milky Way \\citep{Elmegreen1997}.  With new Galactic plane surveys, e.g., the Herschel HiGAL survey \\citep{Molinari2010}, the APEX Telescope Large Area Survey of the Galaxy (ATLASGAL; \\citealt{Schuller2009}), the Bolocam Galactic Plane Survey (BGPS; \\citealt{Aguirre2011}), the H$_2$O Southern Galactic Plane Survey (HOPS; \\citealt{Walsh2011}), and the Millimeter Astronomy Legacy Team 90 GHz Survey (MALT90; \\citealt{Foster2011}), we are on the cusp of better understanding how massive clusters form. The question is, how do we find the massive proto-clusters (MPCs) that will form these YMCs? We investigate how the YMCs may form and provide a simple model with observational properties that can be used to identify MPC candidates. To identify more extreme MPCs in nearby galaxies (e.g., the Antennae Galaxies), we need capable telescopes like the Atacama Large Millimeter/sub-millimeter Array (ALMA).  YMCs may predominantly form from MPCs having gravitational escape speeds greater than the sound speed in photo-ionized gas. When this condition is met, ionization cannot disrupt the entire MPC.  Stars can continue to form from the remaining neutral gas and star formation efficiency (SFE) increases to 30\\% and higher. The remaining mass in the stellar population nullifies the effects of gas expulsion and the cluster will remain bound. If the absolute value of the gravitational potential energy is greater than the expected thermal energy of the plasma in a massive gas clump, and supernovae have not yet occurred, then we consider the object to be an MPC candidate. We compare the effective photo-ionized sound speed of the plasma ($\\Cii$) to the escape velocity ($v_{esc}$) of the clump, which we denote as $\\Omega \\equiv v_{esc} / \\Cii$. A gas clump that has $\\Omega>1$, implying $v_{esc} > \\Cii$, is an MPC candidate while $\\Omega<1$ are not, since gas can be dispersed by the appearance of the first OB stars.  We describe the simple model of massive stellar cluster formation using the $\\Omega$ parameter in \\S\\ \\ref{sec:ymcf}. We make predictions on the MPC's observational properties for Galactic plane surveys and ALMA (extragalactic) in \\S\\ \\ref{sec:predictions}. \\S\\ \\ref{sec:conclusions} discusses the implications of the model and predictions with a summary of the results. ",
        "conclusions": "\\label{sec:conclusions}  We have discussed the possible conditions necessary for YMC formation and how to identify their progenitors, MPCs, in their primordial state regarding their masses, radii, and flux brightness. The key to identifying whether a massive gas clump can form a YMC is the balance between the gravitational potential of the gas clump and the gas kinematics. We characterize this balance as the ratio between $v_{esc}$ at a given radius and the sound speed of the photo-ionized gas, $\\Cii \\sim 10 \\kms$. If $\\Omega = v_{esc} / \\Cii > 1$ (equivalent to, $v_{esc} > \\Cii$) for a clump of gas ($>3 \\times 10^4 \\msun$) then the system will optimally convert the clump to stellar mass and likely form a YMC. We classify such clumps as MPC candidates. If $\\Omega<1$ then the system does not have a deep enough potential well to keep the photo-ionized gas bound, which will lead to rapid gas dispersal and low star formation efficiency. The end product will be a low mass cluster or group that will feed its stars to the field star population over a short time scale. It may be possible that some of these MPC candidates will not form YMCs due to the cruel cradle effect, where the forming cluster is disrupted by nearby massive GMCs \\citep{Kruijssen2011}.  With the HiGAL, ATLASGAL, and BGPS Galactic plane surveys we should be sensitive to the MPCs throughout the Milky Way. The combination of these surveys will provide us a near complete sample of the MPCs in the Milky Way and pave a path for future high-resolution studies. Furthermore, with ALMA's full potential we should be sensitive to MPCs over $>10^7$\\ \\msun\\ in extragalactic systems within 40 Mpc of the Milky Way. This would help us better connect the Galactic and extragalactic MPCs to better understand the formation and evolution of YMCs, some of which could produce long lived ``young globular clusters'' \\citep[see][and references therein]{PortegiesZwart2010}. It is important to note that $r_{vir}$ is the upper limit to these very massive proto-clusters and such objects will likely not be close to such scales. Observing these extragalactic MPC candidates in the nearby galaxies will help constrain the upper limit radii."
    },
    "1208/1208.3158_arXiv.txt": {
        "abstract": "The rapid and irreversible change of photospheric magnetic fields associated with flares has been confirmed by many recent studies. These studies showed that the photospheric magnetic fields respond to coronal field restructuring and turn to a more horizontal state near the magnetic polarity inversion line (PIL) after eruptions. Recent theoretical work has shown that the change in the Lorentz force associated with a magnetic eruption will lead to such a field configuration at the photosphere. The Helioseismic Magnetic Imager has been providing unprecedented full-disk vector magnetograms  covering the rising phase of the solar cycle 24. In this study, we analyze 18 flares in four active regions, with GOES X-ray class ranging from C4.7 to X5.4. We find that there are permanent and rapid changes of magnetic field around the flaring PIL, the most notable of which is the increase of the transverse magnetic field. The changes of fields integrated over the area and the derived change of Lorentz force both show a strong correlation with flare magnitude. It is the first time that such magnetic field changes have been observed even for C-class flares. Furthermore, for seven events with associated CMEs, we use an estimate of the impulse provided by the Lorentz force, plus the observed CME velocity, to estimate the CME mass. We find that if the time scale of the back reaction is short, i.e., in the order of 10 s,  the derived values of CME mass ($\\sim 10^{15}$g) generally agree with those reported in literature. ",
        "introduction": "Solar flares have been understood as an energy release process due to magnetic reconnections in the solar corona \\citep{kop76}. The magnetic fields in the solar corona are anchored in the dense photosphere. Historically, the photospheric magnetic fields were assumed to be unaffected by flares on short time scales because of high mass density there. Of course their long-term evolution is well known to play an important role in  storing the energy and triggering the flares.  \\cite{wang92} and \\cite{wang94} were the first to show observational evidence of flare-related rapid/irrevisible change of photospheric magnetic fields  based on ground-based vector magnetograms.  The most striking but controversial finding at that time was the  increase of magnetic shear along the magnetic polarity inversion line (PIL). Using line-of-sight magnetograms of SOHO/MDI, \\cite{kos01} found that some irreversible variations of magnetic field in the lower solar atmosphere happened very rapidly in the vicinity of PILs at the beginning of the flare of 2000 July 14. \\cite{wan02} analyzed the observed photospheric magnetic flux evolution across 6 X-class flares, and found significant permanent changes associated with all the events. After surveying 15 X-class flares, \\cite{sud05} concluded that the change in the line-of-sight (LOS) magnetic field always occurs during X-class flares.  \\cite{wan06} noticed that when an active region is away from the disk center, the reconnected low-lying fields would cause an apparent increase of the flux in the polarity toward the limb and a decrease in the polarity closer to the disk center.  Until the launch of SDO, these studies were very limited due to the paucity of continuous/consistent high-quality vector magnetogram series. With a nearly continuous coverage over the entire solar disk, vector magnetograms are being obtained from the Helioseismic and Magnetic Imager \\citep[HMI;][]{schou11} on board the Solar Dynamics Observatory (SDO),  making possible extensive studies that achieve a fundamental physical understanding of the observations. A number of recent papers using HMI data have all pointed to the same conclusion that photospheric magnetic fields turn more horizontal immediately after flares and that magnetic shear increases at surface but relaxes in the corona \\citep{wan12,sun12,liuc12}. For example, \\cite{wan12} found a rapid (in about 30 minutes) and irreversible enhancement in the horizontal magnetic field at the flaring magnetic PIL by a magnitude of $\\sim$ 30\\% associated with the X2.2 flare on 2011 February 15. \\cite{petrie12} has analyzed the magnetic field evolution and Lorentz forces in the X2.2 flare on 2011 February 15, and also found a large Lorentz force change coinciding with the eruption.  From the theoretical side, \\cite{hud08} quantitatively assessed the back reaction on the photosphere and solar interior with the coronal field evolution required to release flare energy, and predicted that the magnetic field should become more horizontal after flares. \\cite{wan10} were first to link this idea to observed field changes.  They provided observational evidence of the increase of transverse field at the PIL when vector magnetograms were available. When only the LOS field measurement was available,  they found that if the source active region is not located at the disk center, the measured apparent LOS field changes are consistent with the picture of \\cite{hud08}, i.e., fields turn more horizontal across the PIL. They used the same concept which we mentioned before: due to the projection effect, there is an apparent increase of the flux in the polarity toward the limb and a decrease for the polarity closer to the disk center. More recently, \\cite{fisher12} and \\cite{hud12} further developed analytic modeling, by separately considering Lorentz forces acting on the upper solar atmosphere and the solar interior. The upward momentum of the erupting plasma can be estimated by equating the change in the upward momentum with the Lorentz force impulse acting on the outer solar atmosphere. The authors also argued that the back reaction on the solar interior may be responsible for the sudden sunspot motion on the photosphere and the excitation of seismic waves in the interior.  It is noted that the previous studies were mainly focused on large flares such as X-class or upper M-class events. HMI has been obtaining seeing-free, high-resolution data since 2010 April. In this Letter, we target our study on the magnetic field change associated with flares in a broad range of magnitudes, including C-class events.  We also attempt to find the possible relationship among flare magnitude, field changes, and momentum involved in the eruptions.  In Section 2, we will describe observations and data processing, and will show two examples of case studies.  The statistical results between flare magnitude and field changes will be presented in Section 3, in which we will also discuss a practical method to estimate the CME mass. Section 4 will give the summary and discussion. ",
        "conclusions": "Taking advantage of the newly released HMI vector magnetograms in flare-productive active regions,  we are able to analyze changes of vector magnetic fields associated with 18 flares. This is the first time that such changes are found for small flares down to the GOES C class. The results listed in Tables 1 and 2 agree with previous studies in the following aspects: \\begin{enumerate} \\item All events exhibit a step-like increase of the horizontal magnetic field after flares, with an order of magnitude of \\begin{math} 10^{20}\\end{math} Mx after integrating over the ROI. \\item  The changes are co-temporal with the flare initiation, and the change-over time is about three time bins of the HMI data, i.e., 36 minutes. However, we believe that the reaction time for the field change could be much shorter than this. \\end{enumerate}  From  the statistical studies of the 18 events, we also note the following: \\begin{enumerate} \\item The permanent magnetic field change is always co-spatial with the PIL connecting the two primary flare kernels. \\item Significant linear relationships between the peak GOES X-ray flux and all the following parameters are found: the size of the affected area, the integrated horizontal field change, and the total downward Lorentz force change. \\end{enumerate}  The above findings clearly support the idea of back reaction of surface magnetic fields to the eruption in the corona, as proposed by \\cite{hud08} and \\cite{fisher12}.  The fields are observed to change from a more vertical to a more horizontal configuration. The downward change of Lorentz force reflects such a topological change in magnetic fields. In the photospheric layers, in static equilibrium before and after eruptive events, there should be a balance between the Lorentz force, gas pressure gradients, and gravity. The Lorentz force difference between the post-flare and pre-flare states is the signature of an unbalanced Lorentz force in the solar atmosphere, occurring during the time of the eruption, in which Lorentz forces are balanced primarily by the inertial force of the accelerating plasma.  If the above physics is correct, then the upward CME momentum can be estimated based on the derived impulse associated with the Lorentz force change. We can then estimate the CME mass.  However,  as we already mentioned, an uncertain parameter in the equation is the reaction time associated with the field change. We prefer to use a short time (10 s based on the hard X-ray observation), as the change is observed to occur in a time scale close to the temporal resolution of the HMI data. If a longer time is used, the estimated CME mass will be much larger than the established values in literature. It is  easier to  estimate mass of CMEs for the close-to-limb events based on the white-light image intensity such as that measured by LASCO coronagraph. We are providing an independent method to estimate the CME mass based on the change of the photospheric magnetic field.  This is particularly useful for events closer to the disk center. Our positive correlation between the change of Lorentz force and the peak soft X-ray flux also agrees with the study of \\cite{zhangj04} and \\cite{zhangj06}, in which they found that the CME speed is associated with the soft X-ray flux."
    },
    "1208/1208.3702_arXiv.txt": {
        "abstract": "We have performed numerical simulations of boundary-driven dynamos using a three-dimensional non-linear magnetohydrodynamical model in a spherical shell geometry. A conducting fluid of magnetic Prandtl number $\\Pm=0.01$ is driven into motion by the counter-rotation of the two hemispheric walls. The resulting flow is of von K\\'arm\\'an type, consisting of a layer of zonal velocity close to the outer wall and a secondary meridional circulation. Above a certain forcing threshold, the mean flow is unstable to non-axisymmetric motions within an equatorial belt. For fixed forcing above this threshold, we have studied the dynamo properties of this flow. The presence of a conducting outer wall is essential to the existence of a dynamo at these parameters. We have therefore studied the effect of changing the material parameters of the wall (magnetic permeability, electrical conductivity, and thickness) on the dynamo. In common with previous studies, we find that dynamos are obtained only when either the conductivity or the permeability is sufficiently large. However, we find that the effect of these two parameters on the dynamo process are different and can even compete to the detriment of the dynamo. Our self-consistent approach allow us to analyze in detail the dynamo feedback loop. The dynamos we obtain are typically dominated by an axisymmetric toroidal magnetic field and an axial dipole component. We show that the ability of the outer shear layer to produce a strong toroidal field depends critically on the presence of a conducting outer wall, which shields the fluid from the vacuum outside. The generation of the axisymmetric poloidal field, on the other hand, occurs in the equatorial belt and does not depend on the wall properties. ",
        "introduction": "An electrically conducting fluid driven by viscous forcing exerted at a boundary generates a dynamo if the fluid's magnetic properties -- electrical conductivity and magnetic permeability -- and flow properties can amplify an initial weak magnetic field and ultimately sustain a magnetic field of significant amplitude. This is the most efficient type of forcing to convert the power applied to the system into kinetic energy available for the dynamo, and so is preferred in laboratory experiments designed to study dynamo action, such as liquid metal experiments. In most of these experiments the energy injection scale is the largest scale of the system.  Recently, results from a boundary-driven dynamo experiment, the von K\\'arm\\'an Sodium (VKS) experiment located in Cadarache, France, have shown that the magnetic properties of the boundaries also greatly affect the ability of the flow to maintain a dynamo \\citep{Mon07}. The VKS experiment consists of a cylindrical container filled with liquid sodium, with two counter-rotating impellers at either end. The mechanical forcing exerted by the impellers on the liquid sodium drives a highly turbulent flow. For a sufficiently strong mechanical forcing, dynamo action has been observed that sustains a large-scale magnetic field despite the unconstrained and turbulent nature of the flow. Furthermore, the axisymmetry of the sustained magnetic field (an axial dipole) implies that the turbulent motions are involved in the dynamo process \\citep[\\emph{e.g.}][]{Pet07}. This is an important result in the study of natural dynamos, which operate at very large Reynolds numbers, and mostly produce large-scale magnetic fields. However, dynamo action is only observed in the VKS experiment when the impellers are made of soft iron -- a material with high magnetic permeability, which produces a discontinuity in the magnetic field between the fluid and the impellers. At the highest achievable mechanical forcing in the experiment, dynamo action has never been observed with either stainless steel or copper impellers \\citep{Aum08}. Consequently, elucidating the effect of changes of magnetic permeability of the impellers on the dynamo, and more generally of magnetic boundary conditions, is critical to understanding how the dynamo mechanism operates in the VKS dynamo experiment. This problem is also crucial in other shear-driven systems, such as the plasma Couette experiment in Madison, Wisconsin \\citep{Spe09}, and the spherical Couette liquid sodium experiment in College Park, Maryland \\citep{Zim11}.  The effect of magnetic boundary conditions on dynamo action has been investigated in numerical simulations for both von K\\'arm\\'an type flow (between coaxial rotating disks) and Ponomarenko type flow (cylindrical helical flow) \\citep{Pono73}. Unfortunately, computational limitations prevent numerical simulations from reproducing the same level of turbulence obtained in laboratory experiments. Numerical models include the large-scale mean velocity component, and sometimes smaller scales with  typical viscous length scales much larger than natural or experimental dynamos due to the use of unrealistically high viscosity. Many studies have adopted a kinematic dynamo approach in which the flow is prescribed for all time with no back-reaction from the magnetic field. The imposed mean base flow in these kinematic models can be chosen analytically \\citep{Kai99,Ava03,Gis08,Gis09,Gie10} or based on data from laboratory water experiments \\citep{Mar03,Rav05,Stef06,Lag08}. Some authors adopt a mean-field approach and parametrize the effects of small-scale turbulence through an $\\alpha$-effect, which corresponds to a mean electromotive force that is linear and homogeneous in the large-scale magnetic field \\citep{Ava03,Lag08,Gie10}. A small number of studies use computationally expensive 3D self-consistent models where the velocity is produced by boundary or volume forcing, and the magnetic feedback on the flow is taken into account \\citep{Bay07,Gis08b,Reu11}, but to our knowledge, only \\citet{Rob10} have addressed the problem of magnetic boundary conditions via self-consistent numerical simulations.  All previous numerical studies, using either the mean base flow only, the mean field approach, or 3D self-consistent models, found that enhanced electrical conductivity or magnetic permeability of either the container walls or impellers lead to a reduction of the dynamo threshold measured by a critical magnetic Reynolds number (where the magnetic Reynolds number corresponds to the ratio of magnetic induction over magnetic diffusion). \\citet{Ava03} attribute the reduction of the dynamo threshold to a change in geometry of the electric current lines or the magnetic field lines leading to a reduction of the total ohmic dissipation. \\citet{Gie10} alternatively invoke the reduction of the ``effective'' magnetic diffusivity, that is the magnetic diffusivity averaged over the whole volume of the system, although they acknowledge that this argument does not explain why different magnetic field growth rates are obtained when varying individually either the magnetic permeability or the electrical conductivity of the disks. \\citet{Pet07} argue that the refraction of the magnetic field lines in the soft iron disks (due to the discontinuity of the tangential magnetic field) may act as a shield for the fluid dynamo region between the two disks from the region behind the disks. Indeed, \\citet{Stef06} have shown that the motions of liquid sodium in the region behind the disks is detrimental for the dynamo action in kinematic simulations. \\citet{Rob10} show that finite values of the wall conductance promote dynamo action even when the wall permeability tends to zero. When the wall conductance tends to zero, on the other hand, their model fails to produce a dynamo even for infinite wall permeability. \\citet{Kai99} find that a surrounding wall of the same conductivity as the fluid favors dynamo action up to an optimal thickness. In this case, the ohmic dissipation in the fluid decreases as the electric currents diffuse into the wall. However, they show that thicker walls are detrimental to dynamos that produce time-dependent magnetic field because the skin effect leads to the presence of eddy currents in the wall, and so the ohmic dissipation increases in this case. In an experimental setup similar to VKS but using gallium as working fluid (which has a lower conductivity than sodium and thus a lower magnetic Reynolds number) and applying transverse magnetic fields, \\citet{Ver10} find that the induced axial magnetic field measured in the mid equatorial plane is two to three times larger in magnitude when soft iron disks are used compared to copper or stainless steel disks. They argue that this result (among other observations) is consistent with an induction mechanism in the rotating disks amplified by the distortion of the magnetic field lines by the soft iron.  In general, in boundary-driven system, some essential component of the dynamo likely operates close to the boundaries, so it is perhaps not surprising that changing the magnetic boundary conditions significantly affects the dynamo threshold. Nevertheless, the current physical interpretations of the experimental observations are partly based on assumptions about the flow properties in kinematic dynamo models, and have not been demonstrated in self-consistent models. Moreover, the cylindrical geometry of the VKS experiment makes it difficult to implement realistic boundary conditions numerically, which has led some authors to adopt idealized boundary conditions (\\emph{e.g.} infinite magnetic permeability which implies vanishing tangential magnetic field at the boundary) \\citep{Gis08,Lag08}.  Here we investigate the underlying problem of the role of magnetic boundary conditions in dynamo models through self-consistent three-dimensional magnetohydrodynamical numerical simulations in spherical shell geometry. Our study extends and expands on the work of \\citet*{Rob10} (RGC10 hereafter) who used the boundary forcing exerted by the counter-rotation of the two hemispheric outer walls to drive a mean flow in a spherical cavity. In their study, RGC10 use a thin wall boundary condition which implies that the magnetic field in the outer wall instantly responds to a change of magnetic field in the fluid. However the conditions under which the thin wall limit is appropriate for modeling the experimental setup are unclear. Here, we investigate in more details the role of the outer wall by modeling a wall of finite thickness and finite values for the electrical conductivity and magnetic permeability. The model is self-consistent in the sense that the flow produced by the motions of the rotating boundaries can be adjusted by the Lorentz forces of the sustained magnetic field. Furthermore, no parameterization of the turbulent effects are included in the equations. For fixed forcing and magnetic properties of the fluid, we have examined the effects of varying the properties of the wall (magnetic permeability $\\mu_w$, electrical conductivity $\\sigma_w$, and thickness $h$) on the resultant dynamo. Spherical geometry has the advantage that magnetic boundary conditions can be easily implemented using a toroidal-poloidal decomposition for the magnetic field. The spherical geometry is convenient to study magnetohydrodynamical (MHD) problems numerically but prevents us from studying the exact same flow obtained in the cylindrical VKS experiment. Moreover, the impellers in the VKS experiment consist of flat disks upon which eight curved blades are attached. The effects of the blades on the flow are not reproduced in our numerical setup. Therefore the application of our results to the VKS experiment will remain tentative. ",
        "conclusions": "Through a series of high resolution simulations, we have studied dynamos driven by boundary forcing in a spherical shell geometry, and the effects of varying independently the thickness, electrical conductivity and magnetic permeability of the outer wall.  For an homogeneous system (same magnetic permeability and conductivity in the fluid and the wall) with a magnetic Prandtl number $\\Pm_f=0.01$, the flow is unable to sustain a magnetic field at the forcing used (corresponding to $\\Rey=48193$, about 200 times the critical forcing for hydrodynamical non-axisymmetric instabilities of the base flow). For a wall thickness $h=0.1 r_o$, increasing the wall conductivity, $\\sigma_w$, by a factor 10 or the wall magnetic permeability, $\\mu_w$, by a factor 1000 creates a dynamo. The effects of high $\\sigma_w$ and high $\\mu_w$ are clearly different on the dynamo threshold, so the decrease of the magnetic diffusivity in the wall, \\mbox{$\\eta_w=1/(\\sigma_w \\mu_w)$}, is not the controlling parameter of this problem. The favorable roles of large wall thickness, conductivity and magnetic permeability on dynamo action obtained in our numerical simulations are in agreement with previous numerical studies using different geometry, different flows and in some cases, idealized boundary conditions \\citep{Ava03,Mar03,Rav05,Gis08,Lag08,Gie10}. In particular for a thin wall (thickness $h=0.01 r_o$), we found a good agreement with the results of \\citet{Rob10}, where a similar setup is used with an outer magnetic boundary condition valid in a thin-wall limit.  In our numerical simulations, in both large $\\sigma_w$ and large $\\mu_w$ cases, the dynamo generates a large-scale (axisymmetric) magnetic field. The magnetic field is mostly an axisymmetric toroidal field and an axisymmetric dipolar poloidal component. The axisymmetric toroidal magnetic field, $\\overline{B_T}$, is generated by an $\\omega$ effect, corresponding to the radial shearing of the radial magnetic field in the shear boundary layer located close to the outer wall. The wall plays an essential role in the amplification of $\\overline{B_T}$ in the shear layer. In the large $\\sigma_w$ case, the discontinuity of conductivity allows strong radial gradient of the axisymmetric azimuthal magnetic field, $\\overline{B_{\\phi}}$, or equivalently large latitudinal electric currents in the wall, shielding the induction in the shear layer from the vacuum outside, thereby allowing stronger $\\overline{B_{\\phi}}$ in the fluid. In the large $\\mu_w$ case, the wall forces the poloidal magnetic field to be normal at the fluid-wall interface, imposing strong radial magnetic field across the shear layer and therefore again the generation of stronger $\\overline{B_{\\phi}}$. Similarly to the large $\\sigma_w$ case, a thick wall provides a wide matching region to the vacuum condition thereby allowing a large amplitude of $\\overline{B_T}$ in the shear layer. By filtering the effects of the wall magnetic properties on the different magnetic modes (Section~\\ref{sec:rolewall}), we can reasonably conclude that the essential role of the wall on the dynamo is to allow for large axisymmetric toroidal field $\\overline{B_T}$ in the fluid. The vacuum boundary condition is detrimental for the dynamo by constraining the allowable growth of $\\overline{B_T}$ and so the other magnetic components of the dynamo that feed from it. The presence of a ``shielding'' wall is therefore essential - either thick, of high conductivity or high permeability. We conjecture that without these conditions, even at higher forcings no dynamo will be found if no hydrodynamical bifurcation occurs and for fixed magnetic diffusivity of the fluid. However we emphasize that this argument applies only for shear flows where the $\\omega$ effect occurs adjacent to the outer boundary.  The axisymmetric poloidal magnetic field, $\\overline{B_P}$, is mostly an axial dipole, and is generated in the equatorial belt, where non-axisymmetric motions are strongest. A coherent emf is produced by a narrow range of non-axisymmetric modes with azimuthal symmetry $5\\leq m \\leq 14$. The velocity modes are sheared by the large latitudinal gradients of the zonal flow $\\overline{u_{\\phi}}$ in the equatorial region. This shearing is essential to produce appropriate azimuthal phase shifts between radial and latitudinal velocity components, leading to non-zero azimuthal average of the cross product between velocity and magnetic field of same azimuthal symmetry. This emf has a significant time-average only for a few modes which are selected by the amplitude of the non-axisymmetric velocity, their azimuthal and latitudinal extent and the latitudinal gradients of the zonal flow. Since all these quantities vary with the strength and the geometry of the forcing, we expect different non-axisymmetric dynamo modes to be selected for different forcings.  In the more general context of dynamo theory, we note that this numerical dynamo model operates similarly to the asymptotic dynamo model of \\citet{Bra64} where a large zonal flow generates a strong axisymmetric toroidal magnetic field, and a small deviation of the flow from axisymmetry may be sufficient to produce an axisymmetric poloidal magnetic field to overcome Cowling's theorem.  Having established how our dynamo model operates, it is important to discuss how our results relate to the observations of the VKS experiment. In this work, we use a spherical geometry for numerical convenience. In this geometry and at the large forcings studied here, the shear exerted by the zonal flow is located in a viscous boundary layer at the outer wall, yielding significant influence of the wall parameters on the dynamo action. In the cylindrical von K\\'arm\\'an setup used in the VKS experiment and at Reynolds number of the order of $10^6$, velocity measurements in water show evidence that the largest axial gradients of the shear layer are located in the equatorial mid-plane between the two counter-rotating disks \\citep{Mar03}. Consequently, an ``axial'' $\\omega$ effect, the shearing of the axial magnetic field lines, is thought to be operating in the equatorial mid-plane of the cylinder \\citep{Bou02}. This is a major difference between our numerical work and the VKS experiment. If the shear layer is far from the outer boundary, then the amplification by the wall of the toroidal field may not be operating in the VKS experiment. In the VKS experiment, the generation of the axisymmetric poloidal magnetic field is usually described as the result of an $\\alpha$ effect produced by the helical vortices present between the blades fixed on the rotating flat disks \\citep{Pet07}. In numerical studies, this effect has been parametrized either by adding a source term in the magnetic induction equation of the mean field \\citep{Lag08,Gie10} or by using an analytical formulation of non-axisymmetric flow \\citep{Gis09}. In this scenario, $\\overline{B_P}$ is produced close to the disks, and the generation mechanism is thought to be helped by the high magnetic permeability or electrical conductivity of the disks. In our work, somewhat differently, we found that $\\overline{B_P}$ is produced in the equatorial belt by fluctuating non-axisymmetric motions, and that the wall parameters do not affect this mechanism.  The discrepancy in the location of the active dynamo regions, and therefore the potentially different role of the wall on the different steps of the dynamo feedback loop, may reveal the importance of the geometry of the container or the presence of blades on the rotating walls in these dynamo experiments, both physical and numerical. However the different approaches used to calculate the induction effects (be they self-consistent or parametrized) in the various numerical codes used to simulate the physical experiment may also explain part of the discrepancy. In a forthcoming work, we will study the consequences of the geometry of the forcing on dynamo action, still using a spherical geometry but varying the latitudinal profile of angular velocity of the wall. Here, we find that the latitudinal gradients of the zonal flow play an important role in the selection of the non-axisymmetric dynamo modes sustaining the axial dipole. Forcing confined closer to the poles may simulate the experimental forcing more realistically. A further interesting question is whether there are certain configurations of the forcing which cannot produce dynamo action."
    },
    "1208/1208.4122_arXiv.txt": {
        "abstract": "We present a method for performing Principal Component Analysis (PCA) on noisy datasets with missing values. Estimates of the measurement error are used to weight the input data such that compared to classic PCA, the resulting eigenvectors are more sensitive to the true underlying signal variations rather than being pulled by heteroskedastic measurement noise. Missing data is simply the limiting case of weight=0. The underlying algorithm is a noise weighted Expectation Maximization (EM) PCA, which has additional benefits of implementation speed and flexibility for smoothing eigenvectors to reduce the noise contribution. We present applications of this method on simulated data and QSO spectra from the Sloan Digital Sky Survey. ",
        "introduction": "Principal Component Analysis (PCA) is a powerful and widely used technique to analyze data by forming a custom set of ``principal component'' eigenvectors that are optimized to describe the most data variance with the fewest number of components \\citep{Pearson1901, Hotelling1933, Jolliffe2002}. With the full set of eigenvectors the data may be reproduced exactly, \\ie, PCA is a transformation which can lend insight by identifying which variations in a complex dataset are most significant and how they are correlated. Alternately, since the eigenvectors are optimized and sorted by their ability to describe variance in the data, PCA may be used to simplify a complex dataset into a few eigenvectors plus coefficients, under the approximation that higher-order eigenvectors are predominantly describing fine tuned noise or otherwise less important features of the data. Example applications within astronomy include classifying spectra by fitting them to PCA templates \\citep{Paris2011, ConnollySzalay1999}, describing Hubble Space Telescope point spread function variations \\citep{Jee2007}, and reducing the dimensionality of Cosmic Microwave Background map data prior to analysis \\citep{Bond1995}.  A limitation of classic PCA is that it does not distinguish between variance due to measurement noise {\\it vs}. variance due to genuine underlying signal variations. Even when an estimate of the measurement variance is available, this information is not used when constructing the eigenvectors, {\\it e.g.}, by deweighting noisy data.  A second limitation of classic PCA is the case of missing data. In some applications, certain observations may be missing some variables and the standard formulas for constructing the eigenvectors do not apply.  For example within astronomy, observed spectra do not cover the same rest-frame wavelengths of objects at different redshifts, and some wavelength bins may be masked due to bright sky lines or cosmic ray contamination. Missing data is an extreme case of noisy data, where missing data is equivalent to data with infinite measurement variance.  This work describes a PCA framework which incorporates estimates of measurement variance while solving for the principal components. This optimizes the eigenvectors to describe the true underlying signal variations without being unduly affected by known measurement noise. Code which implements this algorithm is available at \\url{https://github.com/sbailey/empca}~.  \\citet{Jolliffe2002} \\S 13.6 and \\S 14.2 review prior work on PCA with missing data and incorporating weights into PCA. Most prior work focuses on the identification and removal of outlier data, interpolation over missing data, or special cases such as when the weights can be factorized into independent per-observation and per-variable weights. \\citet{GabrielZamir1979}, \\citet{Wentzell1997}, \\citet{TippingBishop1999}, and \\citet{SrebroJaakkola2003} present iterative solutions for the case of general weights, though none of these find the true PCA solution with orthogonal eigenvectors optimally ranked by their ability to describe the data variance. Instead, they find an unsorted set of non-orthogonal vectors which as a set are optimized to describe the data variance, but individually they are not the optimal linear combinations to describe the most variance with the fewest vectors. Their methods are sufficient for weighted lower-rank matrix approximation, but they lack the potential data insight from optimally combining and sorting the eigenvectors to see which components contribute the most variation.  Within the astronomy literature, \\citet{ConnollySzalay1999} discuss how to interpolate over missing data and use PCA eigenspectra to \\emph{fit} noisy and/or missing data, but they do not address the case of how to \\emph{generate} eigenspectra from noisy but non-missing data. \\citet{BlantonRoweis2007} generate template spectra from noisy and missing data using Non-negative Matrix Factorization (NMF). This method is similar to PCA with the constraint that the template spectra are strictly positive, while not requiring the templates to be orthonormal.  \\citet{TsHogg2012} present a more general ``Heteroskedastic Matrix Factorization'' approach to study Sloan Digital Sky Survey spectra while properly accounting for measurement noise. Their underlying goal is similar to this work, though with an algorithmically different implementation.  The methods presented here directly solve for the PCA eigenvectors with an iterative solution based upon Expectation Maximization PCA (EMPCA). \\citet{Roweis1997} describes an unweighted version of EMPCA, including a method for interpolating missing data, but he does not address the issue of deweighting noisy data. We also take advantage of the iterative nature of the solution to bring unique extensions to PCA, such as noise-filtering the eigenvectors during the solution.  The approach taken here is fundamentally pragmatic.  For example, if one is interested in generating eigenvectors to describe 99\\% of the signal variance, it likely doesn't matter if an iterative algorithm has ``only'' converged at the level of $10^{-5}$ even if the machine precision is formally $10^{-15}$.  We discuss some of the limitations of Weighted EMPCA in \\S\\ref{sec:discussion}, but ultimately we find that these issues are not limiting factors for practical applications.  This work was originally developed for PCA analysis of astronomical spectra and examples are given in that context. It should be noted, however, that these methods are generally applicable to any PCA application with noisy and/or missing data --- nothing in the underlying methodology is specific to astronomical spectra. ",
        "conclusions": "We have described a method for performing PCA on noisy data that properly incorporates measurement noise estimates when solving for the eigenvectors and coefficients.  Missing data is simply the limiting case of weight=0. The method uses an iterative solution based upon Expectation Maximization. The resulting eigenvectors are less sensitive to measurement noise and more sensitive to true underlying signal variations.  The algorithm has been demonstrated on toy data and QSO spectra from SDSS. Code which implements this algorithm is available at \\url{https://github.com/sbailey/empca}~."
    },
    "1208/1208.1531_arXiv.txt": {
        "abstract": "The Heterodyne Instrument for the Far Infrared (HIFI) aboard the Herschel Space Observatory has acquired high-resolution broadband molecular spectra of star-forming regions in a wavelength range that is mostly inaccessible from ground-based astronomical observatories.  These spectral surveys provide new insight into the chemical composition and physical properties of molecular clouds.  In this manuscript, we present initial results from the HIFI spectral survey of the Sagittarius B2(N) molecular cloud, which contains spectral features assigned to at least 40 different molecules in a range of physical environments.  While extensive line blending is observed due to the chemical complexity of this region, reliable molecular line identifications can be made, down to the noise floor, due to the large number of transitions detected for each species in the 1.2 THz survey bandwidth.  This allows for the extraction of new weakly emitting species from the line forest.  These HIFI surveys will be an invaluable archival resource for future investigations into interstellar chemistry. ",
        "introduction": "Broadband molecular line surveys are an effective tool for investigating the rich molecular inventory of star-forming regions\\cite{Herbst2009, Turner1991, Nummelin2000, Belloche2008, Belloche2009, Remijan}.  From detailed analyses of the spectral signatures of molecular clouds, their chemical abundances and physical properties (e.g. temperature and density) can be determined.  The derived parameters can then be used to test and guide models of the source and develop our understanding of the physics and chemistry of the interstellar medium.  These surveys are often made available to the public, which allow for other researchers to pursue more detailed investigations or to search for new molecular species in the data.  A number of new radio astronomical observatories, both single-dish telescopes and interferometeric arrays, have recently begun operations that offer significant advances in bandwidth and sensitivity over older facilities.  These new observatories will lead to a dramatic increase in the number of publicly available spectral surveys in the coming years.  Developing strategies to aid in the analysis and interpretation of these spectra, therefore, is an important goal in the field of astrochemistry.  The large number of data channels associated with these surveys necessitates efficient methods for modeling and visualization of the data and results, as the spectrum should be fit globally---characterizing all molecules simultaneously, rather than one at a time, and across the full survey bandwidth.  Millimeter and submillimeter spectra of chemically rich regions are characterized by a dense line forest at lower signal levels, made up largely of features from complex organics, which emit at many frequencies due to their large partition functions.  Many features therefore consist of blends that contain contributions from several molecules.  Once the spectral signatures of known molecules are characterized down to the noise limit, new species, which can be hidden in the spectrum by the features of molecules that emit more strongly, can be identified.  This manuscript presents early results from a spectral survey toward Sagittarius B2(N) acquired using the Heterodyne Instrument for the Far Infrared (HIFI)\\cite{deGraauw2010} aboard the Herschel Space Observatory\\cite{Pilbratt2010}.  Sgr B2(N) has the greatest observed chemical complexity of any molecular cloud in our galaxy, and a large fraction of the $\\sim$170 molecules detected thus far in the interstellar medium\\cite{Woon} were first detected in this source.  This region has been the target of a number of broadband spectral surveys using ground-based observatories in the centimeter and millimeter wavelength regions\\cite{Turner1991, Nummelin2000, Belloche2008, Belloche2009, Remijan}.  HIFI is a high-resolution spectrometer with continuous coverage from 480--1250 GHz and 1410--1910 GHz.  Most of this frequency range is inaccessible from the ground due to atmospheric absorption, so HIFI offers a new look into a well-studied molecular source.  Additionally, by providing a census of which molecules are emitting in each frequency range, this survey can be used to aid in characterizing the submillimeter spectra of other molecular clouds.  We present an analysis of the emission spectrum of methanol in this spectrum, which is detected through a large number of transitions ranging widely in frequency and excitation energy. This data set serves as an example of the Herschel legacy spectral surveys that we expect will be highly useful to the astronomical community in the coming years. ",
        "conclusions": "We anticipate that the Herschel/HIFI spectral surveys collected as part of the HEXOS key program will be a valuable legacy for the molecular astrophysics community.  The survey described here of the Sagittarius B2(N) molecular cloud, in particular, covers a largely unexplored frequency range and provides a near-complete chemical inventory of the most chemically complex source in our galaxy with high sensitivity and spectral resolution.  The large bandwidth of the survey allows for the detection of a large number of lines for each molecular species, which imposes rigorous constraints on the physical parameters of the region.  Additionally, the detailed inventory of which molecules emit in each frequency range in the most chemically rich source in the Galaxy will be a valuable legacy for observers characterizing the submillimeter spectra of other molecular sources.  The approach described in this manuscript points the way forward to methods for the comprehensive analysis of rich spectra.  Unlike previous broadband surveys of Sgr B2(N) and other star-forming regions, where a single band is characterized, the analysis presented here involves modeling the entire 1.2 THz HIFI spectrum consistently.  This is greatly helped by the fact that this spectrum was obtained in a uniform way and on the same instrument, with high calibration and pointing accuracy and without weather variations from band to band due to the stable environment of space.  The frequency coverage of HIFI overlaps with the highest-frequency bands of the ground-based Atacama Large Millimeter/Submillimeter Array (ALMA), which has the potential of measuring spectra in every atmospheric window between 30 and 900 GHz.  The techniques described here for modeling high bandwidth interstellar spectra will therefore be applicable to that very sensitive instrument.  Other facilities will be available in the near future that can measure large bandwidth spectra:  the Stratospheric Observatory for Infrared Astronomy (SOFIA) will have broadband spectral coverage in the far-IR, and existing ground-based millimeter facilities such as the Institut de Radioastronomie Millim\\'{e}trique (IRAM) 30 m telescope are implementing broadband receivers with high spectral resolution, so wide-bandwidth chemical surveys of molecular clouds are becoming increasingly routine.  We also expect that further insight into the physical and chemical conditions of star-forming regions from spectral surveys such as this will be aided by further laboratory and theoretical work from the chemical physics community.  The laboratory spectra of a large number of molecules of interstellar interest have not been recorded, or are incomplete, in the wavelength range of the HIFI instrument, and the rates of a number of relevant molecular processes (such as chemical reactions and collisional excitation) are not well known.  Developments on these fronts will have a critical role in the revolutionary science anticipated from this new generation of radio astronomical instrumentation."
    },
    "1208/1208.0192_arXiv.txt": {
        "abstract": "{ Organic material found in meteorites and interplanetary dust particles is enriched in D and $^{15}$N. This is consistent with the idea that the functional groups carrying these isotopic anomalies, nitriles and amines, were formed by ion-molecule chemistry in the protosolar nebula. Theoretical models of interstellar fractionation at low temperatures predict large enrichments in both D and $^{15}$N and can account for the largest  isotopic enrichments measured in carbonaceous meteorites. However,   more recent measurements have shown that,  in some primitive samples, a large $^{15}$N enrichment does not correlate with one in  D, and that some D-enriched primitive material displays little, if any,  $^{15}$N enrichment. By considering the spin-state  dependence in ion-molecule reactions involving the ortho and para forms of  H$_2$,    we show that ammonia and related molecules can  exhibit such a wide range of fractionation for both $^{15}$N and D in dense cloud cores. We also show that while the nitriles, HCN and HNC, contain the greatest   $^{15}$N enrichment, this is  not expected to correlate with extreme D enrichment. These calculations therefore support the view that Solar System $^{15}$N and D isotopic anomalies have an interstellar heritage. We also  compare our results to existing astronomical observations and briefly discuss future tests of this model. } ",
        "introduction": "The isotopic enrichments measured in the primitive organic matter found in comets, interplanetary dust particles (IDPs) and meteorites  \\citep{Remusat06,Floss06,Aleon10}, probably had their chemical origin in a low-temperature environment that faciliated ion-molecule isotopic exchange reactions, such as the dense interstellar medium (ISM) or the outer protosolar nebula \\citep{MummaCharnley11,Messenger03}.   The enhanced D/H ratios seen in some primitive material \\citep{Messenger00,Robert03,Messenger03,Keller04,Hartogh11} almost certainly originated in low-temperature environments.  Deuterium  enhancement is initiated by the exothermic ion-molecule exchange reaction \\citep{Millar89} going to the right \\begin{equation}\\label{HDReac} { \\rm{H_3^+~~+~~ HD~\\rightleftharpoons~ H_2D^+ ~~+~~H_2~~+~~170~K} } \\end{equation}  In dense cores where CO and other heavy molecules have condensed as ices on to dust grains (i.e. become depleted), H$_2$D$^+$  molecules can continue to react with HD,  producing HD$_2^+$ and D$_3^+$  \\citep{Roberts03}.  This produces very high levels of  deuteration in the remaining gaseous species as well as, through the resulting high atomic D/H ratios, in molecular ice mantles \\citep[e.g.][]{Ceccarelli07}.  A similar  D fractionation chemistry can also occur in the cold  regions of protoplanetary disks \\citep{Willacy07}.  In both Jupiter Family and Oort Cloud comets,  $\\rm ^{14}N/^{15}N $ ratios are measured ($\\sim 130-170$) for both CN and HCN \\citep{Hutsemekers08,BockeleeMorvan08} that are significantly enhanced relative to both the terrestrial and protosolar values \\citep[272 and 440 respectively,][]{Marty12}. Primitive materials can  exhibit large bulk $^{15}$N enhancements ($\\rm ^{14}N/^{15}N \\sim 180-300$), and  high  spatial resolution measurements further show that they also contain distinct `hotspots' which exhibit the greatest D and $^{15}$N enhancements known ($\\rm ^{14}N/^{15}N <100$) \\citep{Messenger00,Aleon03,Floss06,Busemann06,Briani09b,Aleon10}.  Theoretical models of nitrogen isotopic chemistry in molecular cloud cores, where  gaseous molecules are being depleted by accretion on to dust grains, can reproduce the range of $\\rm ^{14}N/^{15}N$ ratios found in primitive matter \\citep[$\\sim 50-280$, ][]{Aleon10} through exchange reactions between $^{15}$N atoms and molecular ions such as $^{14}$N$_2$H$^+$ and HC$^{14}$NH$^+$  \\citep{CharnleyRodgers02,RodgersCharnley08,RodgersCharnley08_ApJ}.   These models predict that N$_2$, NH$_3$, HCN, CN and HNC  molecules should possess the  greatest $^{15}$N enrichments.  There are two $^{15}$N-fractionating pathways \\citep{RodgersCharnley08_ApJ}: a slow one to N$_2$ and ammonia ($\\sim 10^6$ years) and more rapid one to HCN and other nitriles ($\\sim 10^5$ years).  Recent, albeit sparse,  astronomical observations of starless cores indicate that interstellar  ${\\rm  NH_3}$ and ${\\rm  N_2}$ are not as enriched as predicted for the ISM or as measured in the Solar System:  ${\\rm ^{14}NH_3/^{15}NH_3} \\approx 334$ \\citep{Lis10} and $\\rm ^{14}N_2H^+/^{14}N^{15}NH^+   \\approx 446$  \\citep{Bizzocchi10}.   However, the  most pressing  problem for a direct ISM-Solar System isotopic connection concerns the fact  that, while $^{15}$N  and D meteoritic hotspots do seem to correlate in some samples \\citep[e.g.][]{Aleon03}, they clearly do not in others \\citep[e.g.][]{Busemann06,Gourier08}.  Yet, as has been noted by several authors \\citep{Alexander08,Briani09,Bonal10,Aleon10}, the low-temperature interstellar environments most conducive to producing  large  $^{15}$N enrichments, should also produce concomitantly enormous molecular  D/H ratios (e.g. in ammonia or HCN). Thus, one would expect $^{15}$N  and D hotspots to {\\it always} be spatially correlated, contrary to what is seen. This presents a serious challenge for ion-molecule fractionation chemistry.    In cold interstellar environments,  nuclear spin states can be an important factor in determining chemical reaction rates \\citep[e.g.][]{Flower06}. The major interstellar collision partner,   molecular hydrogen,  has a 170.5~K difference in zero point energy between the $para$ (antiparallel spins) and the $ortho$ (aligned nuclear spins) forms;   this allows some endoergic reactions to proceed, even at low temperatures, if a substantial fraction of the H$_2$ is in the higher energy form, $o$-H$_2$. \\citet{Pagani11} demonstrated that a high H$_2$ ortho-to-para ratio (OPR) therefore inhibits (`poisons') the  production of H$_2$D$^+$ in reaction (1), suppressing the overall deuteration.  The fractionation of $^{15}$N may also be influenced by the OPR of H$_2$. Ammonia formation is initiated by the production of N$^+$ from N$_2$ by He$^+$, which then reacts  in \\begin{equation} \\label{amminitReac} {\\rm N^+~+~H_2} ~\\longrightarrow~ {\\rm NH^+~+~H}, \\end{equation} followed by a sequence of ion-molecule reactions with H$_2$ through to NH$_4^+$ and a final  electron dissociative recombination. Reaction \\ref{amminitReac} has an activation energy of $\\lesssim$200 K \\citep{Gerlich93} which, at low temperatures,  can be overcome by the internal energy of \\textit{o}-H$_2$. Thus, ammonia formation and the related $^{15}$N fractionation in dark clouds are dependent on the H$_2$ OPR \\citep{LeBourlot91}, a distinction not made in earlier models \\citep[e.g.][]{RodgersCharnley08_ApJ}.  A recent re-assessment of the OPR dependence in the original experimental data by \\citet{Dislaire12} indicates that  previous work has  overestimated the low-temperature rate coefficient for reaction \\ref{amminitReac}, involving  \\textit{o}-H$_2$,  by almost three orders of magnitude (cf. Le Bourlot 1991).   Thus, the precise abundance of \\textit{o}-H$_2$ could  play a pivotal role in producing a diverse range of D-$^{15}$N fractionation in interstellar precursor molecules.  In this Letter we quantify the effect the H$_2$ OPR has on interstellar $^{15}$N fractionation, and demonstrate that this can account for the isotopic anomalies observed in primitive Solar System materials. ",
        "conclusions": "By considering the ortho-para dependence in ion-molecule reactions involving H$_2$,   ammonia and related molecules are expected to exhibit a wide range of fractionation in both $^{15}$N and D, consistent with those found in primitive Solar System organic matter. These include very large enrichments in both D and  $^{15}$N, as found in previous models, but also low  $^{15}$N enrichment or even depletion, coupled with modest to large D/H values.  Nitrile functional groups on the other hand, while being the most likely source of  $^{15}$N hotspots, are not predicted to correlate with extreme D enrichment.  These calculations support the premise that the  $^{15}$N isotopic anomalies found in meteorites were set in functional groups formed in a cold molecular cloud.  The predicted $\\rm ^{14}N/^{15}N$ and D/H ratios can account for the limited measurements existing for interstellar cloud cores. To further test and constrain these theoretical models, surveys for key molecules and measurements of relevant  $\\rm ^{14}N/^{15}N$  ratios in molecular clouds are required. A more extensive analysis of how correlations between deuterium and $^{15}$N enrichment in specific molecules are affected by spin-state-dependent chemistry is beyond the scope of this Letter and will be considered elsewhere  (Wirstr\\\"om et al. 2012, in preparation). Predictions from these models can be tested by measurements of  various multiply-substituted isotopologues \\citep[e.g. DC$^{15}$N, ${\\rm ^{15}NH_2D }$,][]{Gerin09}, which will be facilitated with the {\\it Atacama Large Millimeter Array} (ALMA)."
    },
    "1208/1208.5853_arXiv.txt": {
        "abstract": "{High-resolution radio observations are ideal for constraining the value of physical parameters in the inner regions of active-galactic-nucleus jets and complement results on multiwavelength (MWL) observations. This study is part of a wider multifrequency campaign targeting the nearby TeV blazar Markarian 421 (z=0.031), with observations in the sub-mm (SMA), optical/IR (GASP), UV/X-ray (Swift, RXTE, MAXI), and $\\gamma$ rays (Fermi-LAT, MAGIC, VERITAS).} {We investigate the jet's morphology and any proper motions, and the time evolution of physical parameters such as flux densities and spectral index. The aim of our wider multifrequency campaign is to try to shed light on questions such as the nature of the radiating particles, the connection between the radio and $\\gamma$-ray emission, the location of the emitting regions and the origin of the flux variability.} {We consider data obtained with the Very Long Baseline Array (VLBA) over twelve epochs (one observation per month from January to December 2011) at 15\\,GHz and 24\\,GHz. We investigate the inner jet structure on parsec scales through the study of model-fit components for each epoch.} {The structure of Mrk 421 is dominated by a compact ($\\sim$0.13 mas) and bright component, with a one-sided jet detected out to $\\sim$10 mas. We identify 5-6 components in the jet that are consistent with being stationary during the 12-month period studied here. Measurements of the spectral index agree with those of other works: they are fairly flat in the core region and steepen along the jet length. Significant flux-density variations are detected for the core component.} {From our results, we draw an overall scenario in which we estimate a viewing angle $2^\\circ < \\theta < 5^\\circ$ and a different jet velocity for the radio and the high-energy emission regions, such that the respective Doppler factors are $\\delta_r \\sim 3$  and $\\delta_{\\rm h.e.} \\sim 14$.} ",
        "introduction": "\\begin{table*} \\caption{Details of the observations.\\label{observations}} \\begin{center} \\tiny \\begin{tabular}{ccccccccc} \\hline Observation & MJD &  \\multicolumn{2}{c}{Map peak} & \\multicolumn{2}{c}{Beam} & \\multicolumn{2}{c}{1$\\sigma$ rms} & Notes  \\\\ date  & & \\multicolumn{2}{c}{(mJy/beam)} & \\multicolumn{2}{c}{(mas $\\times$ mas, $^\\circ$)} & \\multicolumn{2}{c}{(mJy/beam)}  &  \\\\ & & 15\\,GHz & 24\\,GHz & 15\\,GHz & 24\\,GHz& 15\\,GHz & 24\\,GHz&   \\\\ \\hline \\hline 2011/01/14 & 55575 & 348 & 319 &\\ \\ 1.05 $\\times$ 0.65, 15.3  \\ \\ &\\ \\ 0.79 $\\times$ 0.47, 8.84 \\ \\ & 0.19 & 0.18 & No MK, no NL    \\\\ 2011/02/25 & 55617 & 391 & 338 & 1.16 $\\times$ 0.74, 14.4 & 0.64 $\\times$ 0.39, $-$6.52 & 0.35 & 0.18 & NL snowing \\\\ 2011/03/29 & 55649 & 386 & 359 & 1.06 $\\times$ 0.66, $-$5.37 & 0.65 $\\times$ 0.39, $-$4.26 & 0.17 & 0.24 & No HK \\\\ 2011/04/25 & 55675 & 367 & 308 & 0.92 $\\times$ 0.50, $-$3.79 & 0.61 $\\times$ 0.34, $-$3.88 & 0.28 & 0.33 & -  \\\\ 2011/05/31 & 55712 & 355 & 297 & 0.93 $\\times$ 0.51, $-$5.56 & 0.64 $\\times$ 0.35, $-$8.41 & 0.25 & 0.29 & - \\\\ 2011/06/29 & 55741 & 262 & 208 & 0.89 $\\times$ 0.50, $-$7.17 & 0.56 $\\times$ 0.32, $-$12.3 & 0.17 & 0.37 & No LA   \\\\ 2011/07/28 & 55770 & 220 & 197 & 0.91 $\\times$ 0.56, $-$0.89 & 0.60 $\\times$ 0.37, $-$1.14 & 0.20 & 0.30 & -  \\\\ 2011/08/29 & 55802 & 275 & 200 & 0.97 $\\times$ 0.55, 0.26 & 0.63 $\\times$ 0.35, $-$2.70 & 0.18 & 0.27 & No HK  \\\\ 2011/09/28 & 55832 & 264 & 238 & 1.06 $\\times$ 0.67, 16.2 & 0.72 $\\times$ 0.47, 18.4 & 0.26 & 0.24 & No MK  \\\\ 2011/10/29 & 55863 & 261 & 167 & 1.06 $\\times$ 0.69, 0.09 & 0.73 $\\times$ 0.44, $-$3.73 & 0.26 & 0.14 & HK snowing  \\\\ 2011/11/28 & 55893 & 283 & 201 & 1.02 $\\times$ 0.59, 18.0 & 0.70 $\\times$ 0.42, 14.3 & 0.28 & 0.17 & No PT, FD, MK  \\\\ 2011/12/23 & 55918 & 295 & 287 & 0.89 $\\times$ 0.49, 1.82 & 0.61 $\\times$ 0.35, $-$5.74 & 0.15 & 0.21 & No HK  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}  \\begin{figure*} \\centering \\includegraphics[width=7.5cm]{15ghz_stacked.eps} % \\includegraphics[width=7.5cm]{24ghz_stacked.eps} \\caption{Images of Mrk 421 at 15\\,GHz (left panel) and 24\\,GHz (right panel). These two images were obtained by stacking all of the images of the twelve epochs, at the respective frequency. The restoring beam for the 15\\,GHz image is 0.9 mas$\\times$0.55 mas and the peak flux density is 307.2 mJy/beam. For the 24\\,GHz image, the restoring beam is 0.6 mas$\\times$0.35 mas and the peak is 251.9 mJy/beam. The first contour is 0.35 mJy/beam, which corresponds to three times the off-source noise level. Contour levels are drawn at ($-$1, 1, 1.4, 2, 2.8, 4...) in steps of $\\sqrt{2}$.} \\label{maps} \\end{figure*}   Markarian 421 (R.A.=$11^h$\\ $04^m$\\ $27.313943^s$, Dec.=$+38^\\circ$\\ 12\\arcmin\\ 31.79906\\arcsec, J2000) is one of the nearest ($z=0.031$) and brightest BL~Lac objects in the sky. It was the first extragalactic source detected at TeV energies by the Cherenkov telescope at Whipple Observatory \\citep{Punch1992}. The spectral energy distribution (SED) of this object, dominated by non-thermal emission, has two smooth broad components: one at lower energies, from radio band to the soft X-ray domain, and another at higher energies peaking at $\\gamma$-ray energies \\citep{Abdo2011}. The low-frequency peak is certainly due to synchrotron emission from relativistic electrons in the jet interacting with the magnetic field, whereas the high-frequency peak is probably due to the inverse Compton scattering of the same population of relativistic electrons with synchrotron low-energy photons \\citep[synchrotron-self-compton model, SSC, see][]{Abdo2011, Tavecchio2001}. In this framework, multiwavelength (MWL) coordinated campaigns are a fundamental tool for understanding the physical properties of the source, e.g. by studying variability, which is present at all frequencies, but particularly TeV energies where \\citet{Gaidos1996} measured a doubling time of $\\sim15$ minutes. The accurate MWL study and SED modeling performed by \\citet{Abdo2011} revealed some interesting results, such as the size of the emitting region $R$ and the magnetic field $B$, which in the context of the leptonic scenario they constrained to be $R \\lesssim 10^4 R_g$ and $B \\sim$ 0.05 G. However, the details of the physical processes responsible for the observed emission are still poorly constrained. Because of the considerable variability and the broadband spectrum, multiwavelength  long-term observations are required for a good comprehension of the emission mechanisms.  This study is part of a new multi-epoch and multi-instrument campaign, which also involves observations in the sub-mm (SMA), optical/IR (GASP), UV/X-ray (Swift, RXTE, MAXI), and $\\gamma$ rays (Fermi-LAT, MAGIC, VERITAS), as well as at the cm wavelengths with low resolution observations (e.g. F-GAMMA, Medicina). The aim of this observational effort is to shed light on fundamental questions such as the nature of the radiating particles, the connection between the radio and $\\gamma$-ray emission, the location of the emitting regions, and the origin of the flux variability. Very long baseline interferometry (VLBI) plays an important role in addressing these scientific questions because it is the only technique that can resolve (at least partially) the inner structure of the jet. Therefore, cross-correlation studies of Very Long Baseline Array (VLBA) data with data from other energy ranges (in particular $\\gamma$ rays) can provide us with important information about the structure of the jet and the location of the blazar emission.  At radio frequencies, Mrk 421 clearly shows a one-sided jet structure aligned at a small angle with respect to the line of sight \\citep{Giroletti2006}. In this work, we present new VLBA observations to study in detail the inner jet structure on parsec scales. We are able to investigate the evolution of shocks that arise in the jet, by means of the model-fitting technique. In earlier works \\citep{Piner1999, Piner2004}, the jet components show only subluminal apparent motion, which seems to be a common characteristic of TeV blazars. Thanks to accurate measurements of changes on parsec scales, by the VLBA, we can find valid constraints on the geometry and kinematics of the jet.  This paper is structured as follows: in Section 2 we introduce the dataset, in Section 3 we report the results of this work (model fits, flux density variations, apparent speeds, jet sidedness, spectral index), and in Section 4 we discuss results giving our own interpretation in the astrophysical context. We used the following conventions for cosmological parameters:  $H_0=70$ km sec$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.25$ and $\\Omega_\\Lambda=0.75$, in a flat Universe. We defined the spectral index $\\alpha$ such that $S_{\\nu}\\propto\\nu^\\alpha$.  \\begin{figure*} \\centering \\includegraphics[angle=-90, clip, width=0.9\\columnwidth]{gen15.ps}  % \\includegraphics[angle=-90, clip, width=0.9\\columnwidth]{gen24.ps} \\\\ \\caption{Images of Mrk 421 with model fit components for the first epoch at 15\\,GHz (left panel) and at 24\\,GHz (right panel). Levels are drawn at $(-1, 1, 2, 4...) \\times$ the lowest contour, that is at 1.0 mJy/beam for both images, in steps of 2. The restoring beam is shown in the bottom left corner; its size is given in Table~\\ref{observations}.} \\label{components} \\end{figure*} ",
        "conclusions": "The Doppler factor defined as  $$\\delta=\\frac{1}{\\gamma(1-\\beta \\cos\\theta)}$$  \\noindent is a key element in the study of blazars, since it affects various parameters such as the observed brightness, the SED peak frequency, the variability timescale, and more. Modeling of the SED and study of the variability in different wavebands generally require large values of the blazar Doppler factors; in the case of Mrk421, \\citet{Gaidos1996} estimated $\\delta>9$, from the observed TeV variability time of about 30 min and \\citet{Abdo2011} required a Doppler factor between 20 and 50 to reproduce the broadband SED. In turn, VLBI observations can also constrain $\\delta$ by posing limits on $\\beta$ and $\\theta$, as provided by the various arguments discussed in this work.  When closely spaced repeated observations are available, the study of the proper motion is a useful tool in determining the ranges for $\\beta$ and $\\theta$. Surprisingly, several works \\citep[e.g.][]{Giroletti2004b, Piner2008, Piner2010} have reported subluminal motions, sometimes consistent with the component being stationary, in the jets of all the $\\sim$10 TeV blazars for which proper motion studies have been performed in the literature. Thanks to the large number of dual-frequency observations, the fine time-sampling, and the high quality of the data provided by the good $(u,v)$-coverage, we performed a robust identification of the Gaussian components and constrained their motion to be consistent with no displacement at all. At the same time, the high sensitivity and in particular the stacked image place significant constraints on the jet/counter-jet ratio.  The first immediate consequence is that we can reject the hypothesis that the small $\\beta_{\\rm app}$ is solely due to a projection effect, since it would require an unrealistically narrow viewing angle: in the case of component C4b, the upper limit to the observed motion implies a viewing angle $<1.3^\\circ$ to reproduce the observed jet/counter-jet ratio (and even smaller to agree with the high energy limits). If the jets' distribution is isotropic on the sky, the real number of misaligned sources (parent population) is incompatible with these very small values of $\\theta$. For example, in the Bologna Complete Sample selected by \\citet{Giovannini2005} at low frequency (and thus free from Doppler favoritism bias), one would expect fewer than 0.03 sources with $\\theta<1.3^\\circ$ \\citep[see also][]{Tavecchio2005,Henri2006}.  On the other hand, since larger values of the viewing angle capable of reproducing the observed lack of proper motion are incompatible with the jet/counter-jet ratio, we conclude that the pattern velocity cannot be representative of the bulk flow velocity. The Gaussian components obtained in our model fit provide a good description of the visibility data but do not represent well-defined, high-contrast jet features \\citep[see][]{Lyutikov2010}. In our interpretation, the low apparent speeds found imply that the proper motion of Mrk421 does not provide any information about the jet bulk velocity; even on the basis of the sole jet brightness ratio, untenable viewing angles would be necessary to match the pattern and bulk velocities.  What then are the real values of the viewing angle and the jet bulk velocity in Mrk421? To reproduce the observed jet asymmetry, we needed to consider a range of velocities $0.82<\\beta<1$ and angles $0^\\circ<\\theta<35.0^\\circ$. We could exclude the upper range of the $\\theta$ values, since this would not reproduce the high Doppler factors required by high energy observations \\citep{Gaidos1996, Abdo2011}; in particular, we were unable to achieve $\\delta>20, 10, 5, 3$ when $\\theta>3.0^\\circ, 5.7^\\circ, 11.5^\\circ, 19.4^\\circ$, respectively.  Smaller angles thus seem to be favored since they were the only ones consistent with high values of $\\delta$; however, such small angles would still represent a challenge to the observed radio properties. We were able to estimate the intrinsic power of the radio core $P_c^{\\rm intr}$, by debeaming the observed monochromatic luminosity  of the core $P_c^{\\rm obs}$ with the equation  $$P_c^{\\rm obs}=P_c^{\\rm intr} \\times \\delta_c^{2-\\alpha}.$$ With a value of $P_c^{\\rm obs}\\sim6.8\\times10^{23}$ W\\,Hz$^{-1}$ at 15\\,GHz,  $\\alpha=-0.3$, and $\\delta=20$, we obtained for the intrinsic power of the core $P_c^{\\rm intr}\\sim5.8\\times10^{20}$ W\\,Hz$^{-1}$; this value was at the very low end of the typical range of intrinsic power found for different samples of radio galaxies \\cite[e.g.][]{Liuzzo2011}, suggesting that lower values of $\\delta$ provide a more typical core power.  Moreover, the 12 monthly observations have not revealed any dramatic flux-density variability in the core of the source, which further points to a lower Doppler factor for the radio jet. We estimated the variability brightness temperature of the core ($T_{\\rm B,var}$) with the formula proposed by \\citet{Hovatta2009} $$ T_{\\rm B,var}= 1.548 \\times 10^{-32} \\frac{\\Delta S_{\\rm max}d_L^2}{\\nu^2\\tau^2(1+z)},$$ where $\\nu$ is the observed frequency in GHz, $z$ is the redshift, $d_L$ is the luminosity distance in meters, $\\Delta S_{\\rm max}$ is the difference between the maximum value for the core flux density and the minimum value, and $\\tau$ is the variability time. With the values provided by our observations and $\\tau$ $\\sim$ 90 days, we obtaind a value of $T_{\\rm B,var} \\sim 2.1\\times 10^{10}$~K, which does not require any significant beaming.  We similarly calculated $T_{\\rm B}$ for the most compact component following the standard formula \\citep{Piner1999,Tingay1998} $$T_{\\rm B}=1.22 \\times 10^{12} \\frac{S(1+z)}{ab\\nu^2},$$ where $S$ is the flux density of the component measured in Jy, $a$ and $b$ are the full widths at half maximum of the major and minor axes respectively of the component measured in mas, $z$ is the redshift, and $\\nu$ is the observation frequency in GHz. The resulting $T_{\\rm B}$'s are on the order of a few $\\times 10^{11}$ K, only slightly exceeding the limit derived by \\citet{Readhead1994} from equipartition arguments.  Taken together, the lack of superluminal features, the low core dominance, and the weak variability suggest a scenario in which no strong beaming is required in the radio jet. This is not uncommon in TeV blazars \\citep{Piner2004, Piner2008}, but unprecedentedly firm observational support for it has been provided by our intensive campaign. Low values of the Doppler factor, e.g. $\\delta \\sim 3$, can reproduce the observational radio properties, including the jet brightness asymmetry.  We conclude that the Doppler factor must be different in the radio band than the $\\gamma$-ray band. Since we do not expect that the viewing angle changes significantly, this leads us to the necessity of a velocity structure in the jet, as previously discussed by e.g.\\ \\citet{Chiaberge2000},  \\citet{Georganopoulos2003}, and \\citet{Ghisellini2005}. Our images do not provide strong evidence in favour of either a radial or transverse velocity structure, although previous works have revealed a limb brightening in Mrk421, on both a milliarcsecond scale at 43 GHz \\citep{Piner2010} and at $d>10$ mas at 5 GHz \\citep{Giroletti2006}. This would favor the presence of a transverse velocity structure across the jet axis. This structure consists of two components: a fast inner \\textit{spine} and a slower outer \\textit{layer}. Different Doppler factors were obtained depending on whether we measured the speed of the spine or the layer.  A viable scenario for Mrk421 is that the viewing angle is between $2^\\circ$ and $5^\\circ$, which is consistent with the statistical counts of low-power radio sources and the possibility of reaching the high Doppler factors required by SED modeling and high-energy variability. The jet velocity is structured, with a typical Lorentz factor of $\\gamma\\sim 1.8$ in the radio region (yielding $\\delta \\sim 3$), and $\\gamma \\sim \\delta \\sim 20$ in the high-energy emission region. For example, assuming  $\\theta=4^\\circ$, $\\beta_{\\rm h.e.}=0.998$, and $\\beta_{\\rm r}=0.82$, we obtained a value of $\\delta_{\\rm h.e.}=14.3$ and $\\delta_{\\rm r}=3.2$ and successfully reproduced all the observational properties of the source.  In summary, the detailed analysis presented in this paper has largely confirmed with improved quality expectations based on the knowledge so far achieved for TeV blazars. We have also estimated with a good level of significance some important and fundamental parameters ($\\delta$, $\\theta$, $\\beta$, $\\alpha$) that characterize the physical processes in blazars. However, there is still much to be understood and we expect to obtain other significant results from the analysis extended to other wavelengths, particularly in the $\\gamma$-ray domain. Additional works on the dataset presented in this paper are planned, and will deal with, e.g., the 43\\,GHz images and the polarization properties. Moreover, we intend to combine our dataset  with those of other works (e.g.\\ \\citealt{Piner2005} or the MOJAVE survey, \\citealt{Lister2009}) to increase the temporal coverage of the observations and obtain even tighter constraints over a longer time frame. \\vspace{-0.2 cm}"
    },
    "1208/1208.1850_arXiv.txt": {
        "abstract": "The experimental situation of Dark Matter Direct Detection has reached an exciting cross-roads, with potential hints of a discovery of Dark Matter (DM) from the CDMS, CoGeNT, CRESST-II and DAMA experiments in tension with null-results from xenon-based experiments such as XENON100 and LUX. Given the present controversial experimental status, it is important that the analytical method used to search for DM in Direct Detection experiments is both robust and flexible enough to deal with data for which the distinction between signal and background points is difficult, and hence where the choice between setting a limit or defining a discovery region is debatable. In this article we propose a novel  (Bayesian) analytical method, which can be applied to all Direct Detection experiments and  which extracts the maximum amount of information from the data. We apply our method to the XENON100 experiment data as a worked example, and show that firstly our exclusion limit at $90 \\%$ confidence is in agreement with their own for the 225 Live Days data, but is several times stronger for the 100 Live Days data. Secondly we find that, due to the two points at low values of S1 and S2 in the 225 days data-set, our analysis points to either weak consistency with low-mass Dark Matter or the possible presence of an unknown background. Given the null-result from LUX, the latter scenario seems the more plausible. ",
        "introduction": "Despite convincing gravitational evidence for the existence of dark matter (DM) in our Universe (from galactic to cluster scales) its nature remains a mystery. Yet great progress has been made. In particular direct detection experiments have set progressively stronger limits on the properties of dark matter \\cite{Ahmed:2011gh,Akimov:2011tj}, gaining several orders of magnitude in less than a decade for masses in the $10$ GeV to TeV range.  Several direct detection experiments have reported dark matter-like events in their data (e.g. CoGeNT \\cite{Aalseth:2011wp}, CRESST-II \\cite{Angloher:2011uu} and DAMA \\cite{Bernabei:2010mq}), with the most recent positive result coming from the CDMS-Si experiment \\cite{Agnese:2013rvf}. Such hints are in tension with the limits published by the LUX \\cite{Akerib:2013tjd} and XENON100 \\cite{Aprile:2012_new} collaborations. However several authors have claimed that the systematic uncertainties inherent in their analysis may provide a way of reducing such tension \\cite{Davis:2012vy,Savage:2010tg,Hooper:2013cwa}. In addition if one moves beyond the most basic model of DM-quark scattering and considers e.g. inelastic scattering or isospin-violating DM, where the coupling to neutrons and protons is different, then such tension can also be greatly reduced \\cite{Frandsen:2013cna,Frandsen:2011cg,Schwetz:2011xm,Chang:2010yk,Hooper:2011hd}.   Given the present situation, it is essential to exploit all the information contained in the data. In this article we propose a Bayesian approach, based on the information Hamiltonian, with a view to providing the community with a  a novel and robust interpretation of these conflicting experimental signals. This is not the first Bayesian analysis of Direct Detection data \\cite{Arina:2011si}, however our method is distinct in that it extracts the maximum amount of information from the available data, by exploiting the differences between expected signal and background events. For the purpose of illustration, we will make use of data from the XENON100 calibration \\cite{Aprile:2012_new,Aprile:2011hi}. This is an independent analysis of XENON100 data, and will enable us to check and also confront our new method with the collaboration's approach. {This example is also highly relevant for the LUX experiment, which works under a similar principle. }    As we will show for the case where there are signal-like points\\footnote{We define ``signal-like\" data as those consistent with a signal from DM, however we do not wish to make any explicit claim as to their origin, since they may also be consistent with a background interpretation.} in the data our method is particularly powerful, since one can simultaneously set an exclusion limit and define a potential signal region using Bayesian regions of credibility. This is in contrast to current analytical approaches, which usually involve methods designed only to set limits, such as the $p_{\\mathrm{max}}$ method \\cite{Yellin:2002xd}, or the profile Likelihood analysis with the CL$_{\\mathrm{s}}$ method \\cite{Aprile:2011hx}. We do not claim that our method is technically superior for all cases, however our approach is particularly transparent and easily generalised to many different data-sets.    In section \\ref{sec:IH} we first introduce our method and show how to apply it to Direct Detection experimental data in general; this includes a discussion of when to set limits or claim discovery. In section \\ref{sec:x100} we apply our method to data from the XENON100 experiment \\cite{Aprile:2012_new,Aprile:2011hi} as a worked example and conclude in section \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  In this work we have introduced a Bayesian method of analysing data from Dark Matter Direct Detection experiments. Our method takes as input the data itself and the expected signal and background distributions, defined over the whole data-space, which is divided into a grid of two-dimensional pixels. This enables us to take full advantage of the distinct expected distributions signal and background events, and hence to set limits (or discovery regions) without resorting to conservative approximations.     Using data from the XENON100 experiment \\cite{Aprile:2012_new} as a worked example we demonstrated how one would apply our method to Direct Detection data. {This has direct relevance also to LUX experiment \\cite{Akerib:2013tjd}, and any future runs of XENON100.} We have shown that there is merit in looking beyond the $90 \\%$ confidence limit, as hints of signal may be affecting the structure of the Likelihood and Posterior in a non-trivial manner. Indeed an analysis of the XENON100 data from 225 Live Days indicates a weak preference in the data  for a light DM particle. At 50$\\%$ confidence the best fit cross section is in between $6.15 \\cdot 10^{-44} \\, \\mathrm{cm}^2$ and $2.15 \\cdot 10^{-43} \\, \\mathrm{cm}^2$ for an $8 \\, \\mathrm{GeV}$ WIMP; the error bars being relatively large, it is very premature to argue that this is evidence for Dark Matter. Similar regions can be obtained for any dark matter particle with a mass below $\\sim 20$ GeV, with a possible evidence for a dark matter signal in the data vanishing for masses above about 20 GeV. If indeed these points are due to a detection of Dark Matter, more data from the XENON100 experiment should increase the confidence level and shrink the error bars on the cross section. Alternatively, these events may be found to be due to an additional background process or the anomalous component of the ER background, in which case the signal significance would vanish with more data. {Considering the recent null result from the LUX experiment \\cite{Akerib:2013tjd}, the latter would seem to be a more plausible explanation.}    We also demonstrated that our new method can produce a complementary analysis to the one currently used by the XENON100 collaboration, where the data are placed into bands.  Indeed our limit and theirs agree for the most recent 225 Live Days data-set \\cite{Aprile:2012_new}, however ours is several times stronger for the data from 100 Live Days \\cite{Aprile:2011hi}. The reason for this disagreement for the older data-set is not clear. However it is possible that since the background was higher due to krypton contamination, there was a greater proportion of background events leaking into the region where signal was expected (i.e. the more signal-like bands of the analysis used by the XENON100 collaboration), which may have fooled their analysis into setting too weak a limit.  Additionally our method could be even more robust, especially if one exploits the full detector volume (with $f$ and $b$ now depending on physical positions in the detector).    Our analysis can be seen as an independent analysis of the XENON100 data, and more importantly could  be employed by any present or forthcoming experimental collaboration for such a purpose. {In particular, our method can be easily applied to the LUX experiment \\cite{Akerib:2013tjd}, since it operates on a similar principle to XENON100. In this case one should hope to find agreement with our Bayesian results and the frequentist method used by the LUX collaboration, which should provide an important cross-check of the LUX results. Future experiments such as XENON1T \\cite{Aprile:2012zx}, LZ \\cite{Malling:2011va} and SuperCDMS \\cite{Agnese:2013jaa} could also benefit from a Bayesian cross-check.}   The use of our formalism should be very convenient to set limits and potential regions of discovery simultaneously, allowing scenarios where the presence of a signal is ambiguous to be studied without bias. Additionally, our method can be used to go beyond the conservative approach, and to set the strongest limit possible by exploiting the different distributions of signal and background events. With a consistent analytical method used by all dark matter direct detection experiments, the current constraints on the WIMP cross-section should be both stronger and clearer."
    },
    "1208/1208.1311_arXiv.txt": {
        "abstract": "The preliminary results of an analysis of the  KIC~5390438  and KIC~5701829  light curves are presented.  The variations of these stars were detected  by \\citet{baran2011a}  in a search for pulsating M dwarfs in the Kepler public database. The objects have been observed by the Kepler spacecraft during the Q2 and Q3 runs in a short-candence mode (integration time of $\\sim$ 1 min). A Fourier analysis of the time series data has been performed by using the PERIOD04 package. The resulting power spectrum of each star shows a clear excess of power in the frequency range 100 and 350 $\\mu$Hz with a sequence of spaced peaks typical of solar-like oscillations. A rough estimation of the large and small separations has been obtained. Spectroscopic observations secured at the Observatorio Astronomico Nacional in San Pedro M\\'artir allowed us to derive a spectral classification K2III and K0III for KIC~5390438  and KIC~5701829, respectively. Thus, KIC 5390438 and KIC 5701829 have been identified as solar-like oscillating red giant stars. ",
        "introduction": "\\label{introduction}  The Kepler satellite \\citep{borucki}, successfully launched in 2009 March,  is providing  light curves of impressive quality with the primary goal to detect Earth-size planets by means of the transit method. The Kepler spacecraft has onboard  a 0.95-meter diameter telescope with an array of 42 CCDs which is observing continuously a fixed field of view of $\\sim$ 105 square degrees in the constellations Cygnus and Lyra. During the life time of the mission  time series data of more than 150,000 stars will be obtained with a high duty cycle typically from few weeks to several months. Most of these objects will be observed at a near regular long cadence of $\\sim$ 30 minutes, while 512 at a time with a short cadence period of $\\sim$ 1-minute. The high-precision photometry obtained by the satellite is especially well suited for probing the interior of the stars by using the techniques of asteroseismology \\citep{aerts}. This database is unique since it can be used to characterize common oscillation properties of a large number of intrinsic pulsators by analyzing their amplitude spectra \\citep{chaplin, hekker, uytter, baran2011a}  In particular, \\citet{baran2011a} carried out an analysis of a sample of 86 low mass red stars aimed at determine whether radial pulsation are excited to detectable amplitudes in M dwarfs as suggested recently by theoretical predictions \\citep{baran2011b, rodri}. \\citet{baran2011a} derived the spectral types of these 86 objects after which only six of them turned out to be main sequence M stars and none out of six showed signals which could be attributed to M dwarf pulsations. Among these 86 objects a few K-type stars showed short period pulsations which could be attributed to solar-like oscillations. We present here the first results of an analysis of two of these stars: KIC~5390438 and KIC~5701829.    \\begin{figure}[!ht] \\centering \\includegraphics[height=10cm,width=9.0cm]{lfox_fig1.pdf} \\caption{A comparison of the spectrum of KIC~5390438 with two MK standards. From top to bottom: KIC~5390438, K2III standard, K3III standard. The positions of the Mg I triplet and Na doublets absorption features are indicated in each panel.} \\label{fig1} \\end{figure}  \\begin{figure}[!ht] \\centering \\includegraphics[height=10cm,width=9.0cm]{lfox_fig2.pdf} \\caption{A comparison of the spectrum of KIC~5701829 with those of MK standards. From top to bottom: KIC~5701829, K0III standard, K0V standard. The positions of the Mg I triplet and Na doublets absorption features are indicated in each panel.} \\label{fig2} \\end{figure}  \\begin{table}[!ht] \\caption{Observational propierties of the target stars as taken from the Kepler Input Catalogue.\\label{table1}} \\begin{center} \\begin{tabular}{cccccccc} \\hline KEPLER& RA (2000)& DEC (2000)& Kp Mag& $T_{\\rm eff}$& $\\log\\,g$& E(B-V)& Radius\\\\ ID &&&(mag)& (K)& (cm/s$^2$)& (mag)&($R_{\\odot}$)\\\\ \\hline 5390438 & 19 54 32.22 & +40 35 51.0 &10.538& 4536& 4.418& 0.014& 0.924\\\\ 5701829 &19 22 22.64 &+40 59 42.5 &9.047& 4623 &4.634 &0.006 &0.694\\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We have analyzed the light curves of two K-type stars observed by Kepler during aproximately one month. We provided a spectral classification for each star. A preliminary analysis of the power spectrum of the light curves allowed us to get a rough estimate of the large and small frequency separations.  The empirical scaling relations of \\citet{kjeldsen} can be used to provide a first estimate of stellar radius and mass without any stellar modelling:  \\[ \\frac{R}{R_{\\odot}} \\approx \\left (\\frac{135\\,\\mu{\\rm  Hz}}{\\langle \\Delta \\nu \\rangle}\\right)^{2} \\left( \\frac{\\nu_{\\rm max}}{3050\\,\\mu{\\rm Hz}}\\right) \\left (\\frac{T_{\\rm eff}}{5777\\,\\rm{K}}\\right)^{1/2} \\]  \\[ \\frac{M}{M_{\\odot}} \\approx \\left (\\frac{135\\,\\mu{\\rm  Hz}}{ \\langle \\Delta \\nu  \\rangle}\\right)^{4} \\left (\\frac{\\nu_{\\rm max}}{3050\\,\\mu{\\rm Hz}}\\right)^{3} \\left(\\frac{T_{\\rm eff}}{5777\\,\\rm{K}}\\right)^{3/2}\\]   For KIC~5390438, we obtain: $\\frac{M}{M_{\\odot}} \\simeq 1.2$ and  $\\frac{R}{R_{\\odot}} \\simeq 4.0$, while for KIC~5701829, we have $\\frac{M}{M_{\\odot}} \\simeq 1.4$ and  $\\frac{R}{R_{\\odot}} \\simeq 5.8$ . The global asteroseismic parameters derived for KIC~5390438 resemble those of the red giant star KIC~4351319 --a star below the red clump which is still ascending the red-giant branch \\citep{dimauro}.  On the other hand, KIC~5701829 is slightly more massive star in a more advanced evolutionary state than KIC~5390438. These stars will be studied in more details using stellar modelling to retrieve more precise stellar parameters."
    },
    "1208/1208.3314_arXiv.txt": {
        "abstract": "We present modelling of X-ray polarisation spectra emerging from the two competing scenarios that are proposed to explain the broad Fe K$\\alpha$ line in the Seyfert 1 galaxy MCG-6-30-15. The polarisation signature of complex absorption is studied for a partial covering scenario using a clumpy wind and compared to a reflection model based on the lamp-post geometry. The shape of the polarisation percentage and angle as a function of photon energy are found to be distinctly different between the reflection and the absorption case. Relativistic reflection produces significantly stronger polarisation in the 1--10~keV energy band than absorption. The spectrum of the polarisation angle adds additional constraints: in the absorption case it shows a constant shape, whereas the relativistic reflection scenario typically leads to a smooth rotation of the polarisation angle with photon energy. Based on this work, we conclude that a soft X-ray polarimeter on-board a small X-ray satellite may already discriminate between the absorption and the reflection scenario. A promising opportunity may arise with the {\\it X-ray Imaging Polarimetry Explorer (XIPE)} mission, which has been proposed to ESA in response to a small-size (S-class) mission call due for launch in 2017. ",
        "introduction": "For the last two decades, an increasing number of type-1 active galactic nuclei (AGN) showing a broad Fe K$\\alpha$ fluorescent line in the 4--7~keV band has been detected (see e.g. Reeves et al. 2006, Nandra et al. 2007, de la Calle et al. 2010, Patrick et al. 2011). The actual presence of the extended red wing of the line is confirmed in many objects; nonetheless, its physical origin is debated with two major interpretations emerging: a relativistic reflection scenario (Miniutti \\& Fabian 2004) and an absorption scenario (Inoue \\& Matsumoto 2003, Tatum et al. 2012).  The Seyfert galaxy MCG-6-30-15 is an archetypal case among AGN with broad iron lines. Its extended red wing is well-established from long observations with {\\it XMM-Newton} (Wilms et al. 2001, Fabian et al. 2002) and {\\it Suzaku} (Miniutti et al. 2007). Several authors interpret the line as reprocessed X-ray emission emerging from the accretion disc that reaches down to the innermost stable orbit (ISCO) of the supermassive black hole (Miniutti et al. 2003, Reynolds et al. 2009). In this view, the broadening is due to general relativistic and Doppler effects shifting the line centroid as a function of the disc radius. When integrating the emission across the whole disc while taking into account the effects of ray-tracing in a Kerr metric the line is ``relativistically blurred''. Assuming that the accretion disc and its irradiation are indeed truncated at the ISCO, the blurred line puts important constraints on the black hole spin (Fabian et al. 1989, Laor 1991, Dov{\\v c}iak et al. 2004, Brenneman \\& Reynolds 2006).  Following a different approach, Inoue \\& Matsumoto (2003) and Miller, Turner \\& Reeves (2008, 2009) argue that the X-ray data of MCG-6-30-15 can also be explained by assuming several absorbing media located on the line-of-sight and partially covering the primary X-ray source. In this interpretation, the extended red wing is ``carved out'' by absorption and the line shape is much less related to the SMBH spin.  More advanced spectral and timing analyses using forthcoming X-ray missions like {\\it Astro-H} and {\\it NuStar} in addition to {\\it XMM-Newton} may shed more light on how broad iron lines are produced. In this letter, we still explore a different path and test how X-ray polarimetry can help to independently discriminate between the two models.    \\section[]{Comparison of the two scenarios}  The aim of this letter is to have a general view over the two competing scenarios. It is not the scope of this paper to produce an accurate spectral fit to the X-ray data of MCG-6-30-15. We rather assimilate prescriptions for reflection and absorption models that have been presented before and, based on these models, we compute the predicted X-ray polarisation as a function of the observer's viewing angle.  \\begin{figure} \\centering \\includegraphics[trim =30mm 231mm 10mm 6mm, clip, width=13.5cm]{Abs_ref_models.pdf} \\caption{Schematic view of the scenarios considered. $Left$: reflection with a lamp-post geometry and light-bending. $Right$: partial covering with a clumpy wind.} \\label{Fig0}% \\end{figure}   \\subsection{The relativistic reflection model} \\label{sec:reflect}  We first consider relativistic reflection from a cold accretion disc illuminated by an elevated lamp-post on the disc axis. The method is described in detail in Dov{\\v c}iak et al. (2011) so here we only give a brief summary. A grid of local reprocessing models, i.e. taken in the frame of the rotating accretion disc, was computed with the Monte-Carlo radiative transfer code {\\it NOAR} (Dumont, Abrassart \\& Collin 2000) providing the re-emitted intensity as a function of incident and re-emission angle. We defined an isotropic, point-like source emitting an unpolarised spectrum with a power law shape $F_{\\rm *}~\\propto~\\nu^{-\\alpha}$ and $\\alpha = 1.0$. The accretion disc is approximated by a constant density slab with cosmic abundances. Compton scattering, photo-absorption and iron line fluorescent emission are included in the computation of the locally re-emitted spectra. The local polarisation is computed using the transfer equations of Chandrasekhar (1960). Since the reprocessing medium is optically thick, the reprocessing predominately occurs very close to the irradiated surface of the slab and thus the approximation is sufficiently accurate. The local, polarised reflection spectra are then combined with the {\\it KY}-code (Dov{\\v c}iak, Karas \\& Yaqoob 2004) conducting relativistic ray-tracing between the lamp-post, the disc, and the distant observer (see Fig.~\\ref{Fig0}, left). The height of the lamp-post is fixed at 2.5~$R_{\\rm G}$, where $R_{\\rm G}= GM/c^2$, and an extreme Kerr black hole with the dimensionless spin $a = 1$ and a mass of $M = 1.5 \\times 10^6 \\rm M_\\odot$ is assumed.  Our choice of parameters is in good agreement with the assumptions of Miniutti \\& Fabian (2004). We point out though that in our approach the primary source is not off-axis, which should have an impact on the resulting polarisation. Models of a patchy corona (see e.g Galeev, Rosner \\& Vaiana 1979, Haardt et al. 1994) presume that the off-axis sources should be anchored to the disc by magnetic field loops and thus co-rotate in Keplerian motion. In the relativistic reflection model, the X-rays are emitted very close to the black hole and we thus estimate the maximum orbital time-scale occurring in a corona with the radial size $R_{\\rm C} = 50 R_{\\rm G}$ to $317 \\times \\frac{M}{10^7 M_\\odot} \\left[ \\left( \\frac{R_{\\rm C}}{R_{\\rm G}} \\right)^{1.5} + a \\right] [s] \\approx 17$~ks. This time span is by a large factor lower than the minimum exposure time for an observation of MCG-6-30-15 with a near-future X-ray polarimeter (see Sect.~\\ref{sec:xipe}). The observed polarised flux due to a single, off-axis source is thus integrated over many Keplerian orbits. For this reason, the primary emission region should appear axis-symmetric in near-future X-ray polarimetry. So would the irradiation pattern due to a central lamp-post as assumed in our modelling.  The expected X-ray polarisation for a non axis-symmetric, clumpy irradiation pattern of the accretion disc was studied by Schnittman \\& Krolik (2010), who obtain polarisation percentages across the 2--10~keV band that are significantly higher than the results obtained by Dov{\\v c}iak et al. (2011) for the lamp-post geometry. Aside from the different coronal geometry, this is also related to different assumptions about the ionisation of the accretion disk. When assuming a radially structured surface ionisation (Ballantyne et al. 2003) we expect the local percentage of polarisation to increase compared to a neutral disc because the efficiency for electron scattering rises with ionisation. For the same black hole spin, the resulting integrated polarisation observed at infinity must therefore increase as well.  Finally, here we do not include any intrinsic polarisation of the primary radiation. Such polarisation may occur if the primary spectrum emitted by the lamp-post is indeed due to inverse Compton scattering of UV/X-ray photons coming from the disc. This effect may thus strengthen the net polarisation observed at infinity.  In summary, the net polarisation percentage predicted by our lamp-post model with unpolarised primary radiation and a cold accretion disk is likely to be conservative.   \\begin{figure*} \\centering \\includegraphics[trim = 13mm 12mm 47mm 22mm, clip, width=17cm]{PA_PO.pdf} \\caption{Percentage of polarisation $P$ and variation of the polarisation angle $\\Delta\\psi$ with respect to its mean as a function of the energy. Three particular viewing angles $i$ are considered : $30^\\circ$, $45^\\circ$ and $60^\\circ$. $Legend$: a fragmented absorption region (plain line) and a relativistic reflection model with an extreme Kerr SMBH with $a=1$ (red dashed line).} \\label{Fig2}% \\end{figure*}   \\subsection{The complex absorption model} \\label{sec:abso}  An alternative approach to explain the broad red-wing of the Fe~K$\\alpha$ line in MCG-6-30-15 and its lack of variability with respect to the continuum was given by Miller et al. (2008). The authors first suggested a model of absorbed, non-relativistic reflection combined with variable partial covering of the primary source. In the following, Miller et al. (2009) even proposed a pure absorption scenario. This model supposedly is in-line with evidence for high-column density, partial covering absorption found in other AGN (Turner et al. 2009, Reeves et al. 2009, Risaliti et al. 2009). It contains five absorbing zones with the ionised zones 1--3 being required to reproduce the narrow absorption lines in the {\\it Chandra} and {\\it XMM-Newton} grating data. The spectral curvature in the 1--10~keV band is caused by the low-ionisation zones~4 and~5 covering the continuum source by 62\\% and 17\\%, respectively. We therefore focus on these two zones when modelling the expected polarisation. The absorbers 1--3 fully cover the source and they thus represent an additional, low optical depth of $0.0002 < \\tau_{\\rm c} < 0.06$ with respect to Compton scattering that we could add to zones~4 and 5. However, given the large optical depth of zone 4 ($\\tau_{\\rm c} \\sim 1.5$) and its predominant covering factor, it turns out that the impact of zones~1--3 is very limited and that we can safely neglect them.  Using the latest version of the {\\it STOKES} code (Goosmann \\& Gaskell 2007, Marin et al., submitted), we model a geometrically thin, static, disc-like source emitting an isotropic, unpolarised primary spectrum between 1 and 100~keV using the same power law slope of $\\alpha = 1.0$ as for the relativistic modelling presented in Sect. 2.1. This setup is close to the approach of Miller, Turner \\& Reeves (2009). The emitting, central cylinder radially extends up to 0.05~pc and may represent a so-called hot inner flow. Note that in the absorption scenario we do not assume the accretion disc to reach down to the ISCO as otherwise, we would expect again the signs of relativistic reflection. The disk may be truncated at larger radii and thus it presents a low solid angle to the emission region and an eventual reprocessing component remains weak.  Between the source and the observer, a conical, neutral absorber with a height of 1.8~pc along the vertical axis, a half-opening angle of $75^\\circ$, cosmic element abundances, and a Compton optical depth of either $\\tau_{\\rm c} \\sim 1.5$ or $\\tau_{\\rm c} \\sim 0.02$ is defined (Fig.\\ref{Fig0}, right). This parameterisation is a good match to the modelling of zone~4 and zone~5 as given in Miller et al. (2009), except that our computations also include reprocessing. The actual modelling in {\\it STOKES} is done for a uniform density cloud and the clumpiness is included by re-normalising the resulting Stokes fluxes in such a way that 62\\% of the primary emission in the case of zone~4 and 17\\% in the case of zone~5 are incident onto the cloud while 21\\% of the source flux reach the observer directly.    \\subsection{Resulting polarisation signatures}   In Fig.~\\ref{Fig2}, we plot the resulting polarisation as a function of photon energy at a viewing angle of $30^\\circ$, $45^\\circ$ and $60^\\circ$. No circular polarisation can occur in this model setup, so $P$ designates only linear polarisation and $\\Delta\\psi$ the rotation of the polarisation position angle with respect to a convenient average of the polarisation position angles over the depicted energy band. The actual normalisation of the polarisation angle with respect to the disk axis is not of primary interest as we cannot determine it from the observations.  It appears that in comparison with the absorption model, relativistic reflection produces a polarisation percentage $P$ in the 10-50 keV band that is at least by a factor of fifteen higher (Fig.~\\ref{Fig2}, left). At lower energies, $P$ decreases gradually down to 0.1\\% for the reflection scenario while for the absorption model $P$ drops much more drastically below 5 keV. In the relativistic case, the spectral shape of $P$ is determined by the net integration of the polarisation over the accretion disc. The fast motion of the accreting matter and strong gravity effects near the SMBH induce a rotation of the polarisation angle that depends on the position on the disc and on the inclination of the observer. In contrast to this, the energy dependence of $P$ for the absorption scenario is related to the polarisation phase function of electron scattering. A large fraction of the radiation has undergone mostly forward scattering and thus only produces weak polarisation.  Additionally, the primary spectrum of the continuum source favours the emission of soft X-ray photons and thus it causes strong dilution of the transmitted flux by unpolarised radiation at low energies. In the reflection scenario, a significant part of the primary flux is bent down to the disk and the dilution is less efficient.  As the viewing angle increases from $30^\\circ$ to $60^\\circ$, $P$ varies differently in both scenarios. In the relativistic reflection case, $P$ decreases by a factor of $\\sim$ 1.3 as the polarisation contribution from the inner and outer parts of the disc produce differently oriented and partly cancelling polarisation. In the absorption case, $P$ increases with viewing angle; when taken between 5 and 8 keV it is by a factor of $\\sim$ 6 higher at $60^\\circ$ than at $30^\\circ$. Nonetheless, the polarisation percentage for relativistic reflection always remains significantly higher than for the absorption scenario.  The variation of the polarisation angle (Fig.\\ref{Fig2}, right) puts additional constraints on the origin of the broad iron line: in the absorption case, $\\Delta\\psi$ exhibits no variations. The relativistic model, however, induces energy-dependent variations in $\\Delta\\psi$ that increase with viewing angle and that are particularly strong across the iron line. This behaviour is related to the energy dependent albedo and scattering phase function of the disc material. Note that at a viewing angle of $60^\\circ$ the variation of $\\Delta\\psi$ in the 2--10~keV band is larger than $10^\\circ$. At an inclination of $30^\\circ$, which is more probable for MCG-6-30-15, the variation is still around $5^\\circ$. ",
        "conclusions": "The main result so far coming out of our modelling work is that with current observational technology the relativistic scenario should produce measurable soft X-ray polarisation while in the absorption case $P$ should be globally undetectable. If, in addition to that, $\\Delta\\psi$ can be determined to vary across the iron line, a second, independent indicator for the reflection scenario is found.  The exact geometry of the absorber situated along the observer's line-of-sight is unconstrained. To further support the results presented in this paper, we currently explore a range of different absorption scenarios and physical properties of the outflow, and we are going to present their polarisation characteristics in future work. Our modelling of the fragmented medium will be refined by using randomly-situated, spherical absorbers of constant density along the observer's line-of-sight. While being more realistic, such a configuration is expected to produce even lower polarisation percentage than the absorption model presented here. In a clumpy medium, the radiation has to undergo multiple scattering events that have a depolarising effect.  We have also started to investigate the polarisation expected from a wind geometry such as the one suggested by Elvis (2000). In this scheme, the absorbing wind arises vertically from a narrow range of radii on the accretion disc and due to radiation pressure it is bent outward in a conical shape. Similar wind geometries were investigated by Sim et al. (2008,2010) and Schurch et al. (2009) from hydrodynamic simulations. For a distant observer looking at the far-end of the wind, the system is seen in absorption. Preliminary tests show that such wind models produce polarisation that is slightly different from the one obtained for the conical outflows described in Sect.~\\ref{sec:abso} but the results remain within the margins of our conclusions. Note that a partially ionised absorber may produce a stronger reprocessing component than a neutral wind. However, since forward-scattering predominates, the net polarisation is again expected to be low but it may depend on the viewing angle, the geometry of the medium or its ionisation structure. We are going to investigate such scenarios in more detail imposing that they correctly reproduce the observed broad spectral shape of the iron line.  It is also necessary to look into more realisations of the relativistic reflection scenario. In Sect.~\\ref{sec:reflect}, we argue why changes in the irradiation geometry, the ionisation of the accretion disc or the polarisation of the primary radiation with respect to the current model should lead to stronger polarisation. We still need to verify this assumption by adopting the radial ionisation profile used in Svoboda et al. (2012), local reprocessing computations for ionised media and intrinsically polarised X-ray emission.  For now, we summarise our conclusion as follows: if a small, soft X-ray polarimetry mission like {\\it XIPE} observes MCG-6-30-15 and detects polarisation in the 2--10~keV band the reflection scenario is confirmed. If there is no detection of polarisation the absorption scenario is more likely to be correct but a complex reflection model cannot be excluded."
    },
    "1208/1208.4284.txt": {
        "abstract": "The first near-side X-class flare of the Solar Cycle 24 occurred in February 2011 and produced a very strong seismic response in the photosphere. One sunquake was reported by %\\inlinecite{Kosovichev2011}, \\citeauthor{Kosovichev2011} (\\apjl\\ {\\bf 734},  {L15}, \\citeyear{Kosovichev2011}), followed by the discovery of a second sunquake by \\citeauthor{ZGMZ2011} (\\apjl\\ {\\bf 741}, {L35}, \\citeyear{ZGMZ2011}). %\\inlinecite{ZGMZ2011}. The flare had a two-ribbon structure and was associated with a flux rope eruption and a halo coronal mass ejection (CME) as reported in the CACTus catalogue. Following the discovery of the second sunquake and the spatial association of both sources with the locations of the feet of the erupting flux rope (\\citeauthor{ZGMZ2011}\\apjl\\ {\\bf 741}, {L35}, \\citeyear{ZGMZ2011}). %\\cite{ZGMZ2011}, we present here a more detailed analysis of the observed photospheric changes in and around the seismic sources. These sunquakes are quite unusual, taking place early in the impulsive stage of the flare, with the seismic sources showing little hard X-ray (HXR) emission{, and strongest X-ray emission sources  located in the flare ribbons.} %{\\em %The strongest X-ray emission sources are located in the flare loops and so we also consider the photospheric changes there in comparison to the locations of the two seismic sources. } We present a directional time--distance diagram computed for the second source, which clearly shows a ridge corresponding to the travelling acoustic wave packet and find that the quake at the second source happened about 45 seconds to one minute earlier than the first source. Using acoustic holography we report different frequency responses of the two sources. We find strong downflows at both seismic locations and a supersonic horizontal motion at the second site of acoustic wave excitation. ",
        "introduction": "\\label{sec:intro}  Sunquakes are observed as photospheric ripples, which accelerate radially outward from a source region. The theoretical prediction that sunquakes should be produced by the energy released during major solar flares \\cite{wolff72} was supported by their discovery on the Sun by \\inlinecite{kz1998}. The acoustic nature of quakes has been well established since their discovery. However, the exact physical mechanism behind their excitation is still debated with several theories currently under consideration. Observations of sunquakes are relatively rare, possibly due to the difficulties of detecting the photospheric ripples, and helioseismic methods such as time--distance diagram analysis and acoustic holography are employed to look for evidence of acoustic emission. With only a small number of such events verified so far (see \\opencite{Besliu2005}; \\opencite{Donea2006list} for some examples of known quakes from the last solar cycle) the new solar cycle and the virtually continuous high-resolution data  of SDO/HMI \\cite{ScherrerHMI2011,Schou2011hmi} mean that sunquake detections should increase in the coming years.  The 15 February 2011, X2.2 class flare occurred in NOAA active region 11158 and was the first in the much delayed rising activity phase of the new Solar Cycle 24. The active region started emerging in the eastern hemisphere on 10 February 2011, with two bipoles emerging side by side creating a complex multipolar region. As the active region evolved through both emergence and cancellation events, the coronal loops became increasingly sheared resulting in a number of C-class and  M-class flares occurring from 13 February onward, culminating in the X-class event with GOES flux peaking around 01:55\\,UT on 15 February. The X-class flare was a long-duration flare with an impulsive phase, as observed in GOES 1.0 to 0.8 \\AA \\, soft X-ray data, lasting from 01:46 to 01:56\\,UT, and integrated HXR emission observed by {RHESSI} up to $\\approx 100 {\\rm \\ keV}$, peaking just before 01:55\\,UT. The strongest HXR emission was produced in the $6\\,-\\,25{\\rm \\ keV}$ energy band.  \\begin{figure*} \\centerline{ \\begin{tabular}{c @ {\\hspace{1pc}} c} % @ {\\hspace{1pc}} c} \\color{black}{(a)} \\includegraphics[width=.5\\textwidth]{sphys_20110215_eg_sources_nocont_notd.eps}  & \\color{black}{(b)}\\includegraphics[width=.5\\textwidth]{sphys_20110215_ic_sources.eps} \\\\ \\color{black}{(c)} \\includegraphics[width=.5\\textwidth]{sphys_20110215_mg_sources.eps} & %    {\\Large \\bf %    \\hspace{0.0 \\textwidth}  \\color{white}{(a)} %   \\hspace{0.415\\textwidth}  \\color{white}{(b)} %        \\hfill} \\color{black}{(d)}\\includegraphics[width=.5\\textwidth]{sphys_20110215_vv_sources.eps}\\\\ \\color{black}{(e)}\\includegraphics[width=.5\\textwidth]{sphys_20110215_mgsdiff_sources.eps} & \\color{black}{(f)} \\includegraphics[width=.5\\textwidth,bb=50 65 410 280, clip=]{ca11_hx_over_1.eps} %{sphys_20110215_id_boxes} \\\\ \\end{tabular} } \\caption{Sunquake source locations determined from egression power computed at 6 mHz and time--distance overplotted on egression power snapshot (a), continuum intensity (b), magnetogram (c) and velocity (d) images. Panel (e) shows the flare-induced changes in magnetic field computed as the difference between magnetic data in panel (c) and the 20-minute average of HMI line-of-sight magnetorgrams before the flare onset. %The boxes used for mosaic plots are overplotted on intensity difference (f) and magnetic change (e) image at around 01:50 UT. The orange contours in (a)\\,-\\,(e) correspond to 01:51:27 6-mHz egression power contours at 2.5 and 3 times the quiet-Sun egression at this frequency. The red stars in all images mark the locations used for computing time--distance diagrams. Panel (f) is Ca {\\sc ii} K taken by {\\it Hinode}/SOT at  01:51:39\\,UT with HXR contours deduced from RHESSI data in 25\\,-\\,50 keV range overplotted. The HMI data in panels (a)-(e) are Postel projected so the distances along $x$- and $y$-axes are plotted in megameters. Arcsecond coordinates are given along the axes in panel (f). } \\label{fig:srcs_boxes} \\end{figure*}  \\begin{sloppypar} In the photosphere, the flare exhibited a classic two-ribbon pattern, which is most clearly seen in the SDO/HMI line-of-sight magnetic field and velocity running-difference images as well as {\\em Hinode}/SOT Ca {\\sc ii} K observations ({\\it e.g.} Figure \\ref{fig:srcs_boxes}). Though less pronounced, the ribbons are also present in the SDO/HMI continuum data. Spatially, hard X-ray emission was situated  primarily along the ribbons.  The halo CME associated with this flare was detected by CACTus software and is listed in the LASCO catalogue (\\url{http://sidc.oma.be/cactus/catalog/LASCO/2_5_0/qkl/2011/02/}). \\end{sloppypar}  The flare produced a strong seismic response \\cite{Kosovichev2011}, with ripples travelling outward from the source clearly seen in HMI velocity difference data (see, for instance, the online movie in the above article). Using acoustic holography  \\inlinecite{ZGMZ2011} have shown that a second, apparently weaker, source of acoustic waves is present and have shown that the sunquakes occurred at the foot-points of a flux rope. The presence and location of the flux rope has been deduced based on a straight-forward observational case, making use of the photospheric magnetic flux distribution, sigmoidal structure, chromospheric and coronal changes during the eruption. The observational interpretation of the presence of a flux rope is supported by non-linear magnetic-field modelling from HMI vector magnetogram data \\cite{Schrijver2011,XSun2012}. %and Xudong Sun - private communication).  In this article we analyse seismic measurements and report the properties of the detected seismic sources and consider associated changes in photosphere. The data and methods are described in Section \\ref{sec:observation}, results are presented in Section \\ref{sec:results}, followed by discussion and conclusions.   %*******************************************************************************************************************  \\begin{figure*} \\centerline{ \\includegraphics[width=.5\\textwidth]{sphys_td_source11.eps}\\\\ \\includegraphics[width=.5\\textwidth,bb=15 0 500 385, clip=]{sphys_td_source21.eps} } \\caption{Time--distance diagrams for both seismic sources. The first two plots from the left correspond to the first (eastern, strong) source, the following are the western source. Theoretical time--distance curves are overplotted in white. Locations of the sources are marked as red stars in Figure \\ref{fig:srcs_boxes}. } \\label{fig:td_s2} \\end{figure*}  %******************************************************************************************************************* ",
        "conclusions": "We have presented here a time--distance diagram in addition to the one found by \\inlinecite{Kosovichev2011} showing a clear ridge for the second seismic source associated with the 15 February 2011 X-class flare. Using time--distance analysis and HMI line-of-sight velocity observations we deduce that the quakes are excited at around 01:50\\,UT, with the eastern source onset preceding the western one by about $45-60$ seconds. We have also detected apparent horizontal motions of the downward velocity transients at the time and location of both quakes. The speed of such motions is larger than the ambient sound speed. The direction of such motions is  aligned with the stacked egression kernels and amplitude anisotropy of the generated wavefront, indicating that a moving source is the likely scenario for both quakes. We estimate the acoustic energy released by both quakes to be around $1.18\\times 10^{28} {\\rm \\ ergs}$ for Source 1 and $6.08 \\times 10^{27} {\\rm \\ ergs}$ for Source 2. For Source 1 this is about an order of magnitude higher than the Lorentz-force energy estimate for a generic flare provided by \\inlinecite{Hudson08}. This is in line with findings by \\inlinecite{ABMLHC2012}, where more accurate evaluation of Lorentz-force energy has been produced. However, given a number of simplifications used in obtaining the Lorentz-force estimate, such as the use of line-of-sight magnetic field only, the assumption of a single area where changes occur and relatively low magnetic-field strength, in our view it would be premature to discard the Lorentz force as a possible mechanism for quake excitation with further analysis based on \\cite{Fisher2011,Fisher2012} making use of fill vector magnetic-field data necessary.  Further analysis is required in order to understand the physical nature of both detected quakes and their link with the flux-rope eruption that was associated with the X-class flare. In particular, the HMI vector magnetogram data should shed the light on full magnetic changes, and numerical extrapolations of the three-dimensional magnetic field from photosphere through atmosphere and corona will give us a clearer picture of the magnetic-field restructuring and energy release associated with this event."
    },
    "1208/1208.4534_arXiv.txt": {
        "abstract": "{We present new observational constraints on inhomogeneous models based on observables independent of the CMB and large-scale structure. Using Bayesian evidence we find very strong evidence for the homogeneous LCDM model, thus disfavouring inhomogeneous models. Our new constraints are based on quantities independent of the growth of perturbations and rely on cosmic clocks based on atomic physics and on the local density of matter.}   \\begin{document} ",
        "introduction": "The discovery of accelerated expansion about fourteen years ago, based on the dimming of distant supernovae (SN) (\\cite{Riess98, Perl99}), has led to a standard model of cosmology ($\\Lambda$CDM) in which about $72 \\%$ of the energy of the universe is in the form of a cosmological constant (CC) (or dynamical dark energy). An important assumption in this model is that our universe is homogeneous.  As an alternative to a cosmological constant, which has significant theoretical problems due to the required fine-tuning, the supernova data could also be explained without a CC if we live close to the center of a large ($\\gtrsim 1$ Gpc), spherically symmetric void (see \\cite{tomita00, goodwinetal99, celerier00} for some of the first such proposals, while \\cite{mofftat95, mustaphaetal97} discussed this idea even before the measurement of accelerated expansion). The reason is that we cannot distinguish the effect of spatial variations in the geometry from temporal ones in the radial direction from measurements along the past light-cone alone. Specifically, in a large void, the local expansion rate gets larger closer to the center of the underdensity, thus causing the same additional dimming at large distances that would be caused by accelerated expansion in a homogeneous cosmology.  Such void models arguably have larger philosophical and aesthetic problems than the CC. First, the existence of a void of the required size is extremely unlikely in the standard inflationary scenario and its generation would probably involve rather unusual early universe physics. Moreover, there is a large amount of fine tuning associated with the requirement from observation that we find ourselves within $\\sim 1 \\%$ of the void's center (\\cite{alnesamarz06, alnesamarz07, BNV10, blommort10, kodamaetal10, foremanetal10}), a grave violation of the Copernican principle. Additional fine-tuning comes from the fact that the void has to be close to spherical to be consistent with the isotropy that we observe. Finally, the void as an explanation for cosmic acceleration does not actually solve the cosmological constant problem. It merely assumes that the CC is equal to zero without providing any explanation (this criticism of course also applies to most if not all models of dynamical dark energy).  Still, the final say can only come from observation so it is worth comparing predictions of the void model to recent cosmological data. This is not just interesting for the sake of understanding the viability of these specific models, but also serves to test a not often directly tested pillar of our cosmological framework, the assumption of homogeneity (see e.g.~\\cite{clarkmaart10}).   The void is commonly modeled as a Lema{\\^i}tre-Tolman-Bondi universe (LTB, \\cite{Lemaitre97, Tolman34, Bondi47}), which describes the most general spherically symmetric universe filled with pressureless matter (dust). Such models have been extensively used to reproduce the observed supernova magnitudes and can in fact explain {\\it any} Hubble diagram perfectly, provided enough freedom in the void profile is allowed (see e.g.~\\cite{mustaphaetal97, yooetal08}). A more trying test of the LTB cosmology arises when additional data sets are added. Many data sets have by now been considered, including, but not limited to, the primary cosmic microwave background temperature power spectrum (e.g.~\\cite{alnesetal06, ZMS08, Alexetal09, nadatsar11}), spectral distortions of the cosmic microwave background (CMB) (\\cite{caldsteb08}), the kinetic Sunyaev-Zel'dovich (kSZ) effect (\\cite{GBH08b, yooetal10, zibinmoss11, bull12, zhangstebbins11}), the primordial Lithium abundance \\cite{regisclarkson12}, the Baryon Acoustic Oscillation (BAO) scale (e.g.~\\cite{GBH08, ZMS08}), and combinations of the above (\\cite{MZS11, ZMS08, Alexetal09, GBH09, BNV10, ZGBRL12}). In particular, both the combination of the CMB with a low redshift Hubble parameter measurement (and with other data sets) (\\cite{MZS11}), and the observational upper limit on the strength of the kinetic SZ effect in the CMB (\\cite{GBH08b, yooetal10, zibinmoss11, bull12, zhangstebbins11}), appear to rule out at least the simplest versions of the void cosmology.  However, many of the previous studies rely either on the CMB or, through the BAO, on large scale structure (LSS), while it is not universally accepted that these observables are understood well enough in an inhomogeneous universe. Particularly, solving the evolution of density perturbations in LTB is notoriously difficult, although significant progress has by now been made (\\cite{clarksonetal09, Zibin08, alonsoetal10, dunsbyetal10, nishietal12, februaryetal12}). Moreover, constraints derived from the CMB power spectrum may depend on implicit assumptions regarding the distribution of radiation (\\cite{regisclarkson12, clarkreg11}).  It is thus worthwhile considering constraints that do not depend on either CMB or LSS, nor on the details of perturbation evolution in general. For this reason, we will constrain the geometry/expansion of the universe with data complementary to supernovae, by using ``red envelope'' galaxy ages (\\cite{sternetal10, sternetal10b}) as lower bounds on the age of the universe at various redshifts in the range $z = 0 - 1.9$. These estimates of oldest galaxy ages in samples of passively evolving, massive, red galaxies have been used previously to measure the Hubble parameter $H(z)$ as a function of redshift (see \\cite{jimloeb02, simonetal05, sternetal10, sternetal10b}). In fact, this measurement of $H(z)$ has even been used to constrain LTB models in \\cite{FLSC10, wangzhang12}. However, the Hubble determination crucially depends on the assumption that the average formation time of the oldest galaxies in each sample is homogeneous. Since this is not guaranteed in a void model (see Appendix \\ref{app:A}), an issue which was at least discussed in \\cite{wangzhang12}, we cannot safely estimate $H(z)$ from differential age measurements, and use the ages themselves instead. As it turns out, the supernova magnitudes, combined with a local determination of the Hubble rate, prefer a void in which the age of the universe as a function of redshift is low compared to these galaxy ages (see also, e.g., \\cite{ZGBRL12}) so that the addition of the age estimates places strong constraints on the model.  We thus carry out a joint analysis of these three data sets, considering a void that is asymptotically flat and with homogeneous big bang time. We derive parameter constraints and, performing Bayesian model selection, find that $\\Lambda$CDM is significantly favored over LTB models. The data prefer a very low value of the relative matter density in the center of the void ($\\Omega_{\\rm in} \\leq 0.1$), and we show that adding a direct measurement of $\\Omega_m$ from, e.g., clusters causes additional tension between the data sets, helping to further disfavor the LTB model. We thus find results consistent with the studies based on CMB and/or LSS, but using an independent method.  We will briefly review the LTB model and introduce our parametrization in section \\ref{sec:model}, and describe the four types of data used in section \\ref{sec:data}. We then derive parameter constraints, and compare the model to $\\Lambda$CDM using both the $\\chi^2$ statistic and Bayesian evidence comparison, in section \\ref{sec:results}. Finally, we summarize and discuss our results in section \\ref{sec:dis}. The dangers of using red envelope ages to estimate the Hubble parameter in void models are discussed in Appendix \\ref{app:A}. Throughout this article, we will use notation similar to that in \\cite{ZGBRL12} (and previous papers by these authors). ",
        "conclusions": "\\label{sec:dis}  In this work, we have constrained LTB models proposed to explain the apparent accelerated expansion of the universe without introducing a cosmological constant. We have tested the LTB model, together with $\\Lambda$CDM, against several types of data, including a recent compilation of supernova data, and the strongest constraint to date on the local expansion rate. New to this analysis is the addition of age data of old, passively evolving galaxies. We have argued that, while it is not allowed to use the differential ages as a Hubble parameter measurement in the context of void models, the data still have strong constraining power when considered as a set of lower bounds on the age of the universe in a wide redshift range $z \\approx 0 - 1.9$. Finally, we added a cluster-based measurement of the local (relative) matter density $\\Omega_m$ to further constrain the models.   In the previous section, we found that, while in terms of goodness of fit, both the void model and $\\Lambda$CDM provide a satisfactory fit to all data combinations considered, a comparison of the relative merits of each model strongly favors the latter. Specifically, $\\Lambda$CDM provides a better best-fit $\\chi^2$, with the difference ($\\Delta \\chi^2 = \\chi^2_{\\rm void} - \\chi^2_{\\rm \\Lambda CDM}$), after correcting for the fact that the void model has two more free parameters, being $\\Delta \\chi^2 + 2 = 1.6, 5.4$ and $8.0$ for SN + $H_0$, SN + $H_0$ + Ages and SN + $H_0$ + Ages + $\\Omega_m$ respectively. Bayesian analysis provides a good way of comparing the two models against the data through the Bayes factor, which is the ratio of Bayesian evidences. In addition to how well each model is able to fit the data, the advantage of this statistic is that it also takes into account Occam's Razor by penalizing models that need a large parameter space to obtain a good fit. The Bayes factors (here, the ratio of $\\Lambda$CDM over void evidence) are $B =  13.2, 17.0$ and $39.8$ for the data combinations listed above. Therefore, the first two combinations of data sets provide ``strong'' evidence against the void model according to Jeffreys' scale, while the combination of all data even gives ``very strong'' evidence.  Our results are consistent with studies based on the kinetic SZ effect \\cite{GBH08b, yooetal10, zibinmoss11, bull12, zhangstebbins11}, which find that void models that can fit the supernova data are ruled out because they predict a much larger kSZ signal than consistent with observation (although this does assume the matter and radiation perturbations are adiabatic), and with studies combining the CMB temperature spectrum with other data sets. For instance, \\cite{MZS11} shows, among other things, that the CMB in combination with supernovae require a very low value of $H_{00}$, clearly in conflict with the measurement by \\cite{hst11}, and that void models predict a very low clustering amplitude at low redshift, which is also inconsistent with data from clusters and other probes. The novel thing about our work is that it provides independent confirmation of the problems with large void models, without using CMB or LSS data. Moreover, our study suggests that combining the age data we used, with the data used in other studies, including the CMB and LSS data, could lead to significantly larger Bayes factors, even more firmly ruling out the void model. Of course, in that approach one would lose the independence of the details of cosmic perturbations in void models, which was one of the advantages of our approach in this work.  Large values of the parameters governing the void size, $R$ and $\\Delta R$, are not well constrained by our data, with void sizes $\\gtrsim 50$ Gpc still allowed. These parameters have much stronger upper limits when CMB and LSS data are included. For instance, using SN + $H_0$ + BAO + CMB data, \\cite{ZGBRL12} finds $R$ and $\\Delta R$ to be at most a few Gpc, with $R = 0.18^{+0.64}_{-0.18}$ Gpc and $\\Delta R = 2.56^{0.28}_{-0.24}$ Gpc at $68 \\%$ confidence level.   There are some potential caveats to the analysis presented here. First of all, we restricted ourselves to asymptotically flat cosmologies, $\\Omega_{\\rm out} = 1$. In addition, we assumed a homogeneous big bang time such that at early times, the void is just a small perturbation onto a homogeneous Einstein-de Sitter cosmology. Finally, we did not consider the most general void profiles, instead relying on a specific parametrization involving four free parameters. Specifically, we did not include void profiles of the ``unconstrained'' type, which (partially) compensate the central underdensity with an overdensity near the edge of the void.  Our study thus strongly constrains the class of the simplest, perhaps most reasonable LTB models. Relaxing some or all of the above assumptions might improve how well the model fits the data, but also adds more parameter space. It is thus not clear if this would strengthen or weaken the evidence against LTB cosmologies. Our main goal here was to introduce an alternative, strong method to constrain void models, and a study of constraints on the most general void cosmologies is beyond the scope of this article.  In conclusion, void models as an explanation for the apparent cosmic acceleration are under attack from a number of independent types of cosmological data and are by now looking less and less like a viable alternative for a cosmological constant or dark energy."
    },
    "1208/1208.1459.txt": {
        "abstract": "We address a primary question regarding the physical mechanism that triggers the energy release and initiates the onset of eruptions in the magnetar magnetosphere. A self-consistent stationary, axisymmetric model of the magnetar magnetosphere is constructed based on a force-free magnetic field configuration which contains a helically twisted force-free flux rope. The magnetic field configurations in the magnetosphere are obtained as solutions of an inhomogeneous Grad-Shafranov (GS) equation. Given the complex multipolar magnetic fields at the magnetar surface, we also develop a convenient numerical scheme to solve the GS equation. Depending on the surface magnetic field polarity, there exist two kinds of magnetic field configurations, inverse and normal. For these two kinds of configurations, variations of the flux rope equilibrium height in response to gradual surface physical processes, such as flux injections and crust motions, are carefully examined. We find that equilibrium curves contain two branches, one represents a stable equilibrium branch, the other an unstable equilibrium branch. As a result, the evolution of the system shows a catastrophic behavior: when the magnetar surface magnetic field evolves slowly, the height of flux rope would gradually reach a critical value beyond which stable equilibriums can no longer be maintained. Subsequently the flux rope would lose equilibrium and the gradual quasi-static evolution of the magnetar magnetosphere will be replaced by a fast dynamical evolution. In addition to flux injections, the relative motion of active regions would give rise to the catastrophic behavior and lead to magnetic eruptions as well. We propose that a gradual process could lead to a sudden release of magnetosphere energy on a very short dynamical timescale, without being initiated by a sudden fracture in the crust of the magnetar. Some implications of our model are also discussed.   %Specifically, for the inverse configuration, the decay of the %surface magnetic flux would induce the magnetic eruptions, while %for the inverse configuration, the enhancement of surface magnetic %flux would be the trigger. %According to our calculations, w ",
        "introduction": "%This paper is motivated by Low \\& Smith (1993) and Komissarov %(2006). Weinberg (1976) Two intimately connected classes of young neutron stars $-$ Soft Gamma-ray Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs), which are commonly referred to as magnetars, both show high energy emissions (Mazets et al. 1979; Mereghetti \\& Stella 1995; Kouveliotou et al. 1998; Gavriil et al. 2002). It is widely believed that the X-ray luminosity in these sources is powered by the dissipation of non-potential (current-carrying) magnetic fields in the ultra-strongly magnetized magnetosphere with the magnetic field $\\mathbf{B}$ $\\sim 10^{14}-10^{15}$G (Duncan \\& Thompson 1992; Thompson \\& Duncan 1996; Thompson et al. 2002). Occasionally, a much brighter outburst has been observed, i.e., a giant flare releases a total energy of $\\sim 10^{46}$ergs and has a peak luminosity of $\\sim 10^{44 - 46}$ergs $\\mathrm{s}^{-1}$ (for recent reviews see Woods \\& Thompson 2006; Mereghetti 2008). %%%Highly tangled magnetic fields are created %%%inside the neutron stars during supernova core collapse through %%%the dynamo processes (Duncan \\& Thompson 1992, Thompson \\& Murray %%%2001). The interior twisted current-carrying magnetic field exerts %%%the Lorentz force on the crust. When the stress induced by the %%%Lorentz force exceeds the critical stress of the crustal lattice %%%strain, the crustal motion would occur. The resulting crustal %%%motion would transport the current (or twist) inside the neutron %%%star into the magnetosphere. The dissipation of the magnetospheric %%%current supports the persistent emission and a sudden %%%reconfiguration of the magnetic field may explain the flares %%%(Thompson, Lyutikov \\& Kulkarni 2002; Lyutikov 2006). Although the energy for magnetar outbursts is widely believed to be supplied by the star's magnetic field, the physical process by which the energy is stored and released remains one of the great puzzles in high-energy astrophysics. Two possibilities exist for the location where the magnetic energy is stored prior to an eruption: in the magnetar crust or in the magnetosphere. For the former possibility, a giant flare may be caused by a sudden untwisting of the internal (to the neutron star) magnetic field (Thompson \\& Duncan 2001). The subsequent quick and brittle fracture of the crust leads to energetic outbursts\\footnote{However, recent calculations by Levin \\& Lyutikov (2012) imply that plastic deformations of the crust are more likely to occur and the crust model of giant flares may not explain the fast dynamical energy release.}. During the outbursts, there would be an enhanced twist of the magnetospheric magnetic field lines. In this crust scenario, the energy stored in the external twist is limited by the tensile strength of the crust. Alternatively, due to the difficulties to explain the short timescale of the giant flare rise time, $\\sim$ 0.25 ms (Palmer et al. 2005), the second possibility --- the magnetospheric storage model, was proposed by Lyutikov (2006). In this particular scenario, the energy stored in the external twist need not be limited by the tensile strength of the crust, but instead by the total external magnetic field energy (Yu 2011b). In the magnetospheric storage model, the magnetic energy storage processes take place quasi-statically on a longer timescale than the dynamical flare timescale prior to the eruption. %%A mechanism of this type would easily account for the fact that %%large eruptions typically occur on the Alfven timescale in the %%magnetosphere.  %%another unresolved important aspect for this model, i.e., %%how the magnetosphere lose the equilibrium and erupts for the %%magnetar magnetosphere. In the magnetospheric model for giant flares, the energy released during an eruption is built up gradually in the magnetosphere before the eruption. Some interesting properties about the storage of magnetic energy of the magnetospheric models have been discussed in Yu (2011b). But there still remains a primary question regarding magnetospheric models, i.e., %which is to identify what is the mechanism that triggers the energy release and initiates the eruption? More specifically, how a very gradual process by the flux injections (Klu\\'{z}niak \\& Ruderman 1998; Thompson et al. 2002) or crust motions (Ruderman 1991) could lead to a sudden release of magnetosphere energy on a very short dynamical timescale, without being initiated by a sudden fracture in the rigid component of the neutron star. %However, there still remains a major problem, i.e., how a very %gradual processes could lead to the sudden release of %magnetosphere energy on a very short timescale. This catastrophic behavior is essentially reminiscent of solar flares and coronal mass ejections (CMEs). It is conceivable that the magnetosphere adjusts quasi-statically in response to the slowly-changing boundary conditions at the magnetar surface. After reaching a critical point, the magnetosphere could no longer maintain a stable equilibrium and a sudden reconfiguration of the magnetic field occurs due to loss of equilibrium (Forbes \\& Isenberg 1991; Isenberg et al. 1993; Forbes \\& Priest 1995). The subsequent physical processes would proceed on a dynamical time scale. This catastrophic process naturally explains the puzzle how a very slow process could lead to the sudden release of external magnetic energy on a much shorter timescale (Thompson et al. 2002). %%%and also some recent work on the magnetar crust model (Levin \\& %%%Lyutikov 2012), which sheds doubts on the crust origin of the %%%magnetar giant flares, %%%we propose a magnetosphere model for magnetar giant flares.   %%In this paper we aim to address this problem. %%or certain kind of dynamical instability (Lyutikov 2003; %%Komissarov et al. 2007) produces the giant flare, in close analogy %%with .The energetics of magnetar giant flares are considered %%carefully in our previous work (Yu 2011b). In this paper, we focus %%on another important aspect of giant flares. %%Crust model and magnetosphere model for giant flares. %%For simplicity, we ignore GR effects in this paper.  %%Large-scale eruptive CMEs often give rise to the opening up of %%magnetic field lines that were originally closed. The processes of %%magnetic fields opening up have been extensively investigated in %%the CME studies (Barnes \\& Sturrock 1972; Aly 1984; Mikic \\& %%Linker 1994). The similarity between solar eruptions and magnetar %%giant flares (Lyutikov 2003) motivates this study on the magnetic %%energy buildup process in the magnetar atmosphere. %%We note that there are important differences between solar %%eruptions and magnetar outbursts.  %%It is worthwhile to note %%that these two mechanisms for the flare eruption are expected to %%share many similarities since they both involve dissipation of %%magnetic energy in the magnetosphere. %%General relativistic effects are currently, however, not taken %%into account in relevant energy storage processes. In this work we %%will investigate these processes with GR spacetime curvature %%effects considered. More specifically, %%and describe the background geometry of the magnetar atmosphere %%with the Schwarzschild metric. %%Schwarzschild metric around the central neutron %%star. , we will ignore the rotation effects and %%would like to explore the virial theorem with %%spacetime curvature considered and  %{\\bf flux rope discussion} %{\\it I may put a schematic fig here to illustrate the basic %mechanism of flux rope formation. (Shibata's 2009 ApJ paper).} %The magnetospheric eruption model for magnetar giant flares is %quite similar to solar eruptions. The magnetar giant flares may involve a sudden loss of equilibrium in the magnetosphere, in close analogy to solar flares and CMEs (Lyutikov 2006). % This paper aims to answer this question. % page 345 in Thompson's paper 2002. A number of CMEs show structures consistent with the ejection of a magnetic flux rope\\footnote{The flux rope is a helically twisted magnetic arcade anchored on the solar surface and often used to model prominences in the solar corona.}, as has been reported by Chen et al.(1997) and Dere et al. (1999). Hence magnetic flux ropes have been presumed to be typical structures in the solar corona, and their eruptions might be closely related to solar flares and CMEs (Forbes \\& Isenberg 1991; Isenberg et al. 1993). Similarly, in the magnetar magnetosphere, magnetic flux ropes could be generated due to the pre-flare activity (G\\\"{o}tz et al. 2007; Gill \\& Heyl 2010). As the magnetic flux injects from deep inside the magnetar, the dissipation of the magnetic field may give rise to the precursor activity. The magnetic dissipation of the precursor could also lead to topology changes of the magnetic fields and the formation of a magnetic flux rope\\footnote{In this work, we do not address the question of how a flux rope might be formed. Possible mechanism had been discussed by van Ballegooijen \\& Martens (1989).}. Such a flux rope is also an indispensable ingredient for the radio afterglow observed in SGR1806 (Gaensler et al. 2005; Lyutikov 2006). It is worthwhile to note that the magnetic field interior to the flux rope, which is suspended in the magnetosphere, is helically twisted. It corresponds to a locally twisted feature in the magnetosphere (Thompson et al 2002; Pavan et al. 2009). Such locally twisted flux ropes seem to be more consistent with recent observations, which suggest the presence of localized twist, rather than global twist (e.g. Woods et al. 2007; Perna \\& Gotthelf 2008).  %{\\bf multipolar discussion} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Observations show a striking feature that the emergence of a strong four-peaked pattern in the light curve of the 1998 August 27 event from SGR 1900+14, which was shown in data from the Ulysses and BeppoSAX gamma-ray detectors (Feroci et al. 2001). %The pulse profile of the 1998 August 27 giant flare provides hints %for the presence of multipolar magnetic fields in SGR 1900+14: %four subpulses of a large amplitude appeared during the %intermediate portion of the burst (Feroci et al. 2001; Thompson \\& %Duncan 2001). These remarkable data may imply that the geometry of the magnetic field was quite complicated in regions close to the star. It is reasonable to infer that, near the magnetar surface, the magnetic field geometry of an SGR/AXP source involves higher multi-poles. %As a result, complex boundary conditions should be important for %the outburst of magnetars. The multipolar magnetic field configurations could be readily understood within the magnetar model. Physically speaking, the electric currents, formed during the birth of magnetar, slowly push out from within the magnetar and generate active regions on the magnetar surface. These active regions manifest themselves as the multipolar regions on the magnetar surface. Due to the presence of the active regions, the magnetic field may deviates from a simple dipole configuration near the magnetar surface (Pavan et al. 2009). Our calculations show that multipolar magnetic active regions, especially their relative motions, would have important implications for the catastrophic eruptions of magnetar giant flares (see Section 4). %%%Prior observations may indicate that multipolar structure of %%%magnetosphere plays an important role in the magnetar outbursts. %%%In this paper we would investigate how the complex multipolar %%%field affect catastrophic behavior of the magnetar magnetospheres. %%%In the solar corona, the majority of prominences are formed formed %%%between active regions (Tang 1987 Solar Physics, 107, 233). This %%%suggests that the multipolar structure of the corona fields play %%%an important role in the prominence eruptions. The dynamical %%%process of flux rope formation in multipolar active regions has %%%been numerically investigated by (van Ballegooijen \\& Martens %%%1989; Mackay \\& van Ballegooijen 2005).  Motivated by the similarity between giant flares and solar CMEs, Lyutikov (2006) speculated that magnetar giant flares may also be trigged by the loss of equilibrium of a magnetic field containing a twisted flux rope. But no solid calculations about the equilibrium loss of a flux rope in magnetar magnetosphere have been performed yet. In this paper we focus on the possibility of magnetospheric origin for giant flares and propose that the gradual variations at the magnetar surface could lead to fast dynamical processes in the magnetosphere. We will construct a force-free magnetosphere model with a flux rope suspended in the magnetosphere and study the catastrophic behavior of the flux rope in a background multi-polar magnetic field configuration, taking into account the possible effects of flux injections (Klu\\'{z}niak \\& Ruderman 1998; Thompson et al. 2002) and crust horizontal motions (Ruderman 1991). We are especially interested in the critical height of the flux rope that can be achieved in our model. %complex multipolar magnetospheric magnetic field. %Due to the complex multipolar boundary conditions at the magnetar surface, %the usually adopted Green function method is hard to use. Instead, In the mean time, we also develop a convenient numerical scheme to solve the inhomogeneous Grad-Shafranov (GS) equation. Since observed magnetars have a very slow rotation rate, we ignore rotation effects throughout this work.  %%%It is physically reasonable to assume that the pre-flare closed %%%state must possess more magnetic energy than the post-flare open %%%state. As will be discussed in detail below, requiring the %%%magnetic flux to open imposes an extreme energy constraint on %%%theories for CMEs. This energy requirement on solar CMEs has been %%%under extensive theoretical studies in the past decades (Aly 1984; %%%Sturrock 1991; Wolfson \\& Dlamini 1997; Zhang \\& Low 2005). The %%%energy storage processes take place quasi-statically on a long %%%timescale. When the magnetic field reaches a threshold, due to the %%%instability or loss of confinement, the field erupts suddenly on a %%%much shorter dynamical timescale. Analogous processes of magnetic %%%field opening up are believed to occur in magnetar giant flares %%%(Woods et al. 2001; Thompson et al. 2002; Beloborodov 2009). All %%%these features of the storage model are in good agreement with the %%%observational behavior of magnetar giant flares  %%%One drawback of prior analytic treatment is that the radius of the %%%flux rope, $r_0$, is supposed to be small enough compared to the %%%height of the flux rope, $h$ (Lin et al. 1998). So in this paper %%%we choose to attack this problem numerically and avoid the %%%shortcomings of the analytic treatment. The method developed in %%%this paper thus has a wider applications than an analytic %%%analysis.  Not correct statement??  %%%The formation before the coronal mass ejection has been %%%extensively discussed by van Ballegooijen \\& Martens (1989), %%%Mackay \\& van Ballegooijen (2006) and Green \\& Kliem (2009). %%%Self-similar multipolar arcade is discussed by Lynden-Bell's %%%student at cambridge and Volpi et al. (2009) in MNRAS.  This paper is structured as follows: in Section 2 we describe the basic equations for the force-free magnetosphere model as well as the multipolar boundary conditions. Two possible magnetic configurations are also discussed in this section. In Section 3 we will discuss the internal and external equilibrium constraints in our model. Numerical results about catastrophic behaviors of the magnetosphere in response to flux injections and crust motions are discussed in Section 4. Conclusions and discussions are given in Section 5. Technical details about the force-free magnetosphere magnetic field are given in Appendix A and B. ",
        "conclusions": "\\label{sec:diss} In this work, we consider the possibility that the primary mechanism for driving an eruption in magnetar giant flare is a catastrophic loss of equilibrium of a helically twisted flux rope in the magnetar magnetosphere. The loss of equilibrium behavior of a flux rope is investigated in a multi-polar magnetic field configuration, taking into account possible effects of flux injections and crust horizontal motions. The loss of equilibrium model describes a quasi-static equilibrium that varies in response to slow changes at the magnetar surface. Beyond a critical point, the stable equilibrium can not be maintained and the transition to a dynamical evolution naturally occurs.  Equilibrium states of a stationary, axisymmetric magnetic field in the non-rotating magnetosphere containing a flux rope are obtained as solutions of the inhomogeneous GS equation in a spherical polar coordinate. In view of the complex multipolar boundary conditions at the magnetar surface, we develop a numerical method to solve the GS equation. Two kinds of magnetic field configurations, inverse and normal, are carefully examined in this work. Both of them present the loss of equilibrium behavior. We carefully examined the critical height of the flux rope beyond which a stable equilibrium could not be maintained and a sudden release of magnetosphere will be triggered. We find that the critical flux rope height is different for the two types of configurations. We also investigate effects of another form of boundary changes, crust horizontal motions, on the loss of equilibrium behavior. Our results show that both the flux injection and crust motions could trigger the catastrophic behavior in the magnetosphere.    %The equilibrium states of axisymmetric force-free magnetic fields %in magnetar magnetospheres are found as solutions of the %Grad-Shafranov equations. We study two kinds of boundary behaviors %that possibly occur at the magnetar surface.  %Motivated by crust model failure and complex multipolar magnetic %field at the surface, we construct general relativistic models of %non-rotating neutron stars endowed with strong magnetic field.  %Current sheet is a good topic for further investigations (Wolfson %1985) and its effects on the catastrophic behavior is interesting %to study.  %A newly derived general relativistic magnetic virial theorem in %presented in this work. Based on this magnetic virial theorem, we %carefully examine the GR effects on the well known Aly-Sturrock %energy threshold. We found this energy threshold decreases with %the magnetar mass. As a result, more massive magnetars are more %prone to eruptions. The non-force-free magnetic field induced by %the mass loading is further investigated as a possibility to %bypass the Aly-Sturrock constraint for typical magnetar mass $\\sim %1.4 M_{\\odot}$.  In our simplified model, the flux rope is assumed to be a closed current ring encircling the magnetar. It is suspended in the magnetar magnetosphere and two ends of the flux rope are not anchored to the magnetar surfaces. %In our simplified model, the %ends of the flux rope are not anchored to the neutron star %surface. We expect that the overall catastrophic behavior of our model should remain the same even the anchoring effects of flux rope is taken into account. In order to further understand the anchoring effects on the catastrophic behavior, more realistic three dimensional model that includes a flux rope with two ends anchored to the magnetar surface is worth further investigation. %a three dimensional model needs to be developed.  The magnetic energy that could be released in our model is an interesting issue that is worthwhile to be explored further. %%%%%Though these works were performed in a Cartesian coordinate, these %%%%%results may still have guidance. For a purely dipolar boundary %%%%%condition, previous studies have shown that about 1\\% (Forbes \\& %%%%%Isenberg) However, these calculations are performed in a Cartesian %%%%%coordinate.  so cartesian may not carry over to our model. For a spherical coordinate in our paper, the ground energy reference state, based on which the fraction of released magnetic energy can be calculated, is the fully open Aly-Sturrock field (Aly 1984, 1991; Sturrock 1991; Yu 2011b). Since the boundary conditions in our paper is more complex (not dipolar or quadrupolar boundary conditions), the construction of the Aly-Sturrock field is technically nontrivial. As a result, the fraction of the energy that could be released in our model involves a very careful calculation of the Aly-Sturrock field. Intuitively, according to prior studies, which have shown that magnetic configurations with more complex boundary conditions would be able to release more energies (Forbes \\& Isenberg 1991; Isenberg et al. 1993), as well as the complex boundary conditions adopted in our model, it is expected that our model would release enough energy for a magnetar giant flare\\footnote{Typically 1\\% of the magnetic energy release could already account for a giant flare. For a simple dipolar boundary condition, about 1\\% of the magnetic energy can be released (Forbes \\& Isenberg 1991). However, about 5\\% of the magnetic energy can be released for a quadrupolar boundary condition (Isenberg et al. 1993).}. Full details of the energetics of our model would be discussed elsewhere.  It is possible that the current sheet forms after the system loses equilibrium (Forbes \\& Isenberg 1991; Forbes \\& Priest 1995). With the formation of current sheet, the tearing instability would develop inside the current sheet (Komissarov et al. 2007) and the subsequent magnetic reconnection would further accelerate the flux rope (Priest \\& Forbes 2000). %is conducive to the further acceleration of the flux rope via the %magnetic reconnection processes . Magnetic field configurations with the current sheet in a spherical polar coordinate is a long-standing unresolved problem. We note that numerical method developed in this work can be further extended to allow the presence of current sheets (Yu 2011b). Further discussions about the current sheets formation and their effects on the catastrophic behavior would be reported in a separate paper.  It would be interesting to to examine the spectral properties of the model presented in this paper (Thompson et al 2002), %calculate the emission spectrum of the model shown in this paper, in which a locally twisted flux rope is self-consistently incorporated into the magnetar magnetosphere. By fitting the spectral features with observations (Pavan et al. 2009), certain parameters in this model, e.g., flux rope height, electric current and magnetic field, may be better constrained. In parallel, recent Fermi observation of Crab nebulae gamma-ray flare could possibly be explained by the magnetic reconnection models (Abdo et al. 2011), in which the loss of equilibrium may be the trigger for the formation of current sheet and subsequent magnetic reconnection processes. The high energy flare emission from Crab Nebulae is thought to be synchrotron radiation by relativistic electron-positron pairs accelerated in this current sheet (Uzdensky et al. 2011). It would be instructive to calculate, based on the model shown in this paper, the emission spectra and compare them with recent Crab Nebula observations.  %Current sheet formation is an interesting issue for further %studies. And Crab nebulae super-flare on April 12 can be probably %explained by the same model as magnetar giant flares,  they are %caused by sudden rearrangements of the magnetic field not far from %the neutron star. Super-flares occur as the intense magnetic field %near the pulsar undergoes sudden restructuring. Such changes can %accelerate particles like electrons to rapid velocities near the %speed of light. As these high-speed electrons interact with the %magnetic field, they emit powerful gamma rays.  And it is %interesting to explore the spectrum properties of the model %presented in this paper.  The models constructed in this work are likely to be useful as initial states in high resolution force-free electrodynamic numerical simulations to explore the dynamics of magnetic eruptions (Yu 2011a). Our current model can not address the dissipation processes that occur during giant flares. This is left for a future work to directly simulate the behavior of loss of equilibrium and relevant dissipation processes using a newly developed resistive force-free electrodynamic code (Yu 2011a). %In this work, we have not address the flux rope formation problem. %Detailed three dimensional numerical modelling of the the %formation and evolution of the flux rope is an interesting work to %be done in future (Yu 2011a).      %The field topology change from a closed state to an open state %must be accompanied by the magnetic reconnection. After the loss %of confinement of the toroidal magnetic field, the gradual %quasi-static evolution of the magnetar's magnetosphere will be %replaced by the dynamical evolution of the field. The magnetic %energy dissipation in the strongly magnetized plasma is caused by %the tearing mode instaiblity (Lyutikov 2003, Komissarov et al. %2007). Relativistic reconnection in the nonlinear regime needs %further studies to better understand the magnetar outburst %behaviors.   %One good way to distinguish crust model and magnetosphere model is %the simultaneous radio and X-ray observation of a giant flare. The %reconnection type event may be accompanied by coherent radio %emission resembling solar type III radio bursts."
    },
    "1208/1208.0584_arXiv.txt": {
        "abstract": "We report the discovery of a one-sided 3.6\\arcsec\\ (24 kpc, projected) long jet in the high-redshift, $z$=4.72, quasar GB~1428+4217 in new \\Chandra\\ X-ray and VLA radio observations. This is the highest redshift kiloparsec-scale X-ray/radio jet known. Analysis of archival VLBI 2.3 and 8.6 GHz data reveal a faint one-sided jet extending out to $\\sim$200 parsecs and aligned to within $\\sim$30\\deg\\ of the \\Chandra/VLA emission. The 3.6\\arcsec\\ distant knot is not detected in an archival \\HST\\ image, and its broad-band spectral energy distribution is consistent with an origin from inverse Compton scattering of cosmic microwave background photons for the X-rays.  Assuming also equipartition between the radiating particles and magnetic field, the implied jet Lorentz factor is $\\approx 5$. This is similar to the other two known $z \\sim 4$ kpc-scale X-ray jet cases and smaller than typically inferred in lower-redshift cases. Although there are still but a few such very high-redshift quasar X-ray jets known, for an inverse Compton origin, the present data suggest that they are less relativistic on large-scales than their lower-redshift counterparts. ",
        "introduction": "}  The $z=4.72$ quasar \\gb\\ (B3~1428+422) was identified in a search for high-redshift objects through targeted optical spectroscopy of flat-spectrum radio sources \\citep{hoo98,fab97}. It is a luminous X-ray source with detected X-ray and radio variability characteristic of a blazar \\citep{fab99,wor06,ver10}. On parsec-scale, VLBA 15 GHz images indicate a dominant high brightness temperature, $T_{\\rm b} \\simeq (4-6) \\times 10^{11}$ K, core component with a faint one-sided jet-like extension \\citep[][and references therein]{ver10}. The high luminosity, variability, and radio compactness are all properties consistent with Doppler beaming of emission from a relativistic jet aligned close to our line of sight \\citep[see][]{fab99}.  Such high-redshift radio/X-ray sources offer a unique glimpse into powerful outbursts from active galactic nuclei (AGN) in the early Universe. As the ambient medium into which large-scale jets propagate (including host galaxy environments and intergalactic medium) is expected to be drastically different at such early epochs \\citep[e.g.,][]{dey06,mil08}, studies of large-scale jet structures allow us to probe radio source interactions with their environment. Motivated by X-ray detections of kpc-scale jets in two very high-redshift quasars, 1745+624 at $z=3.9$ \\citep{che06} and GB~1508+5714 at $z=4.3$ \\citep{sie03,yua03}, and the noticeable dearth of \\Chandra\\ observations of z$\\simgt$2 jet systems \\citep[e.g.,][]{kat05,har06,mas11}, we began a program to obtain arcsecond-resolution radio and X-ray imaging of more such systems \\citep{che05,che08} with the aim to understand the physics of the highest-redshift relativistic jets.  Using \\Chandra\\ X-ray Observatory and NRAO\\footnote{The National Radio Astronomy Observatory is operated by Associated Universities, Inc. under a cooperative agreement with the National Science Foundation.} Very Large Array (VLA) imaging observations of \\gb, we discovered an X-ray/radio jet separated from the nucleus by 3.6\\arcsec\\ (24 kpc, projected)\\footnote{Adopting $H_{\\rm 0}=71~$km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.27$ and $\\Omega_{\\rm \\Lambda}=0.73$, the quasar is at a luminosity distance, $D_{\\rm L}=44.5$ Gpc and 1\\arcsec = 6.59 kpc.}. At $z=4.72$, this is the most distant kpc-scale jet imaged in X-rays. No significant optical emission is detected from the 3.6\\arcsec\\ knot in an archival {\\textit{Hubble Space Telescope (HST)}} image. To probe smaller scale emission, we imaged archival very long baseline interferometry (VLBI) 2.3 and 8.6 GHz data, revealing a one-sided $\\sim$200 parsec long jet aligned within $\\sim 30\\deg$ of the kpc-scale \\Chandra/VLA structure. In the following, we present these multi-wavelength observations (Section~2), and go on to discuss the physical parameters of the large-scale outflow in terms of inverse Compton emission models, comparing this case to other X-ray detected AGN jets (Section~3). ",
        "conclusions": "The kpc-scale radio and X-ray jet reported here in the $z=4.72$ blazar \\gb\\ is the highest-redshift example thus far. Together with the blazar nature of its core emission (Section~1), the one sidedness of the kpc-scale and $\\sim$200 pc scale jet seen in our VLBI images imply that the jet is relativistic and probably aligned at a small angle, $\\theta$, to our line of sight. For $\\theta \\simlt 20\\deg$, the 3.6\\arcsec\\ knot distance from the core corresponds to 24/(sin~$\\theta$) kpc $\\simgt 70$ kpc, deprojected for the detected jet. This is far enough from its parent host galaxy that the most significant source of seed photons for inverse Compton scattering is the cosmic microwave background (CMB). This is especially relevant at $z$=4.72, where the CMB energy density for an observer at rest at the source redshift is $u^{*}_{\\rm CMB} = 4.2 \\times 10^{-13}~(1+z)^{4} = 4.5 \\times 10^{-10}$ erg cm$^{-3}$, i.e., 1070 times greater than it is locally.  The radio and X-ray morphology of \\gb\\ at $z=4.72$ is nearly identical to that of the $z=4.3$ quasar GB~1508+5714 with similar X-ray \\citep{sie03,yua03} and radio \\citep{che04,che05} luminosities in the jet. Both show single radio features separated from the bright nucleus and display steep spectra (with spectral indices, $\\alpha_{\\rm r} = 1.0 \\pm 0.1$ and $1.4 \\pm 0.2$, respectively). Because of poor statistics, the X-ray spectrum in \\gb\\ is unconstrained and can not be compared with that of GB~1508+5714 where $\\alpha_{\\rm x} \\approx 0.9 \\pm 0.4$ was determined. In \\gb, we calculate the ratio \\fxfr\\ $\\equiv (\\nu_{\\rm x} F_{\\rm x})$ / $(\\nu_{\\rm r} F_{\\rm r})$ = 205, using monochromatic flux densities at 1 keV and 1.4 GHz. This is larger than the $z=4.3$ case where \\fxfr\\ = 158, and larger than found in other lower-$z$ examples \\citep[cf.,][]{che04,mas11}. Equivalently, \\fxfr\\ = 205 corresponds to a radio to X-ray spectral index, $\\alpha_{\\rm rx} = 0.72$. Although $\\alpha_{\\rm rx}$ is just consistent with the limit on the radio to optical spectral index, $\\alpha_{\\rm ro} > 0.73$, it is smaller than the radio spectral index, $\\alpha_{\\rm r} = 1.0$, indicating that the X-rays are not a simple extension of the radio synchrotron spectrum.  The X-ray emission could correspond to either inverse Compton emission of the low-energy electrons involving CMB target photon field \\citep[`IC/CMB' model;][]{tav00,cel01}, or an additional synchrotron component due to a higher-energy electron population \\citep[e.g.,][]{sta04,hard06}. Overall, the broad-band spectral indices appear consistent with the (1+$z$)$^{4}$ amplification of the CMB energy density, \\fxfr\\ $\\sim (\\delta/\\Gamma)^{2}\\, u'_{\\rm CMB}/u'_{\\rm B} \\propto (1+z)^{4}~(\\delta/B)^{2}$ (with the Doppler factor, $\\delta$, the bulk Lorentz factor, $\\Gamma$, and the CMB and magnetic field energy densities in the jet rest frame, $u'_{\\rm CMB} = \\Gamma^{2} \\, u^{*}_{\\rm CMB}$ and $u'_{\\rm B} = B^{2}/8\\pi$, respectively), as would be expected in the IC/CMB model \\citep[e.g.,][]{sch02,che04}. Following \\citet{che04}, we apply the IC/CMB model assuming also equipartition with a relativistic electron spectrum with power-law slope, $p = 2 \\alpha + 1 = 3$, extending down to a minimum $\\gamma =10$, in a sphere with an apparent size of $1.8 \\times 1\\arcsec$ for the \\gb\\ jet knot \\citep[see][]{mar83}. For the case where the jet Lorentz factor is set equal to the Doppler factor, we derive $\\Gamma = \\delta = 4.7$, and a magnetic field in the jet rest frame, $B=35$ $\\mu$G. These parameters are fairly insensitive to the extrapolation of the $p=3$ slope down to low energies and our conservative upper limit to the radio knot size. Assuming instead a typical observed radio jet spectral index of 0.75 at lower-redshifts \\citep[following][]{kat05}, and the knot diameter is $0.6\\times$ smaller than the assumed 1\\arcsec\\ limit, we obtain $B=21$ $\\mu$G and $\\delta = 6.3$. These estimates assume a single X-ray/radio emitting zone applies to our $\\sim 1\\arcsec$ (6.6 kpc) angular resolution element.  Similarly applying the IC/CMB model to the other high-redshift X-ray jet, GB~1508+4714 \\citep[$z=4.3$;][]{che04} and 1745+624 \\citep[$z=3.9$;][]{che06}, they derived $\\delta \\sim 3-5$. These values are smaller than the typical ones ($\\delta \\sim 5-10$) determined for lower-redshift jets \\citep[e.g.,][]{kat05}. Although the number of quasar X-ray jets found at such very high-redshifts ($z \\sim 4 - 5$) is still small, the slower large-scale jets implied in our analysis for the current sample seem to indicate that they are either: (1) intrinsically less relativistic, or (2) decelerate more rapidly out to $\\sim$ 10's $-$ 100's kpc-scales than their lower-redshift counterparts. In the first case, this would be consistent with the recent finding that the relative abundance of high-power blazars relative to the parent population of radio-loud quasars and radio galaxies decreased substantially above $z \\sim 3$, likely as a result of a decrease of the average bulk Lorentz factor of blazar jets in the early Universe \\citep[see][]{vol11}. The latter possibility could be supported by the fact that the environments of high-redshift ($z > 2$) radio sources are believed to reside in more inhomogeneous and multi-phase environments than in nearby radio sources \\citep[see e.g.,][]{ree89,dey06}, and this may manifest in the slower large-scale jets due to an enhanced entrainment of the ambient gas when the jet propagates through the host galaxy.  An argument put forth against the IC/CMB model comes from observations of the unresolved (at \\Chandra\\ resolution) quasar ``cores'' at high-redshift.  In this interpretation for the kpc-scale emission, one may expect that the X-ray emission in these cores could also be enhanced, as they contain an unresolved portion of the jet (which we know are highly relativistic from, e.g., VLBI superluminal motion studies).  This could manifest in different X-ray core properties at high-$z$ (X-ray, and optical-to-X-ray spectra), but \\Chandra\\ studies thus far show no such trends with redshift \\citep[e.g.,][]{lop06}. However, these analyses are hampered by the fact that most of the sources analyzed so far in this context are only moderately radio-loud, and in general, only a small fraction of jets in radio-loud quasars extend to kpc-scales \\citep[e.g.,][]{bri84,liu02}.  The radio structures of higher-redshift samples are now being investigated in more detail \\citep[][Gobeille et al., in preparation]{gob11}, and will pave the way for further \\Chandra\\ X-ray studies. X-ray observations of a larger sample of kpc-scale radio jets over a broad redshift range ($z>2$) should provide further elucidation of the correct X-ray emission mechanism which will give us insight on the deeper physics contained in the data.  Such high redshift systems allow us to study jets which result from the ``earliest'' actively accreting black hole systems. As most detections are currently at $z\\simlt 2$, we have begun to study more distant examples and these results will be presented in a future paper."
    },
    "1208/1208.0298_arXiv.txt": {
        "abstract": "The Cosmic Microwave Background (CMB) is a relict of the early universe. Its perfect 2.725K blackbody spectrum demonstrates that the universe underwent a hot, ionized early phase; its anisotropy (about 80 $\\mu$K rms) provides strong evidence for the presence of photon-matter oscillations in the primeval plasma, shaping the initial phase of the formation of structures; its polarization state (about 3 $\\mu$K rms), and in particular its rotational component (less than 0.1 $\\mu$K rms) might allow to study the inflation process in the very early universe, and the physics of extremely high energies, impossible to reach with accelerators. The CMB is observed by means of microwave and mm-wave telescopes, and its measurements drove the development of ultra-sensitive bolometric detectors, sophisticated modulators, and advanced cryogenic and space technologies. Here we focus on the new frontiers of CMB research: the precision measurements of its linear polarization state, at large and intermediate angular scales, and the measurement of the inverse-Compton effect of CMB photons crossing clusters of Galaxies. In this framework, we will describe the formidable experimental challenges faced by ground-based, near-space and space experiments, using large arrays of detectors. We will show that sensitivity and mapping speed improvement obtained with these arrays must be accompanied by a corresponding reduction of systematic effects (especially for CMB polarimeters), and by improved knowledge of foreground emission, to fully exploit the huge scientific potential of these missions. ",
        "introduction": "\\label{sec:intro}  %  We live in an expanding universe, cooling down from a state of extremely high density and temperature, the big bang. In our universe the ratio between the density of photons (the photons of the cosmic microwave background) and the density of baryons is of the order of 10$^9$: this abundance of photons dominated the dynamics of the Universe in the initial phase (first 50000 years). During the first 380000 years the universe was ionized and opaque to radiation, due to the tight coupling between photons and charged baryons. Radiation thermalized in this primeval fireball, producing a blackbody spectrum. When the universe cooled down below 3000K, neutral atoms formed (recombination), and radiation decoupled from matter, traveling basically without any further interaction all the way to our telescopes. Due to the expansion of the universe, the wavelengths of photons expand (by the same amount all lengths expanded, a factor of 1100). What was a glowing 3000K blackbody 380000 years after the big bang, has been redshifted to millimeter waves, and is now observable as a faint background of microwaves. This is the cosmic microwave background, which has been observed as a 2.725K blackbody filling the present universe \\cite{Math99}.  The CMB is remarkably isotropic. However, it is widely believed that the large scale structure of the universe observed today (see e.g. \\cite{Padm93}) derives from the growth of initial density seeds, already visible as small anisotropies in the maps of the Cosmic Microwave Background. This scenario works only if there is dark (i.e. not interacting electromagnetically) matter, already clumped at the epoch of CMB decoupling, gravitationally inducing anisotropy in the CMB. There are three physical processes converting the density perturbations present at recombination into {\\it observable } CMB temperature fluctuations $\\Delta T / T$. They are: the photon density fluctuations $\\delta_\\gamma$, which can be related to the matter density fluctuations $\\Delta \\rho$ once a specific class of perturbations is specified; the gravitational redshift of photons scattered in an over-density or an under-density  with gravitational potential difference $\\Delta \\phi_r$; the Doppler effect produced by the proper motion with velocity $v$ of the electrons scattering the CMB photons. In formulas: \\begin{equation} \\label{dtt} \\frac{\\Delta T}{T}(\\vec{n}) \\approx \\frac{1}{4}\\delta_{\\gamma r}+ \\frac{1}{3} {\\Delta \\phi_r \\over c^2} - \\vec{n} \\cdot {\\vec{v_r} \\over c} \\end{equation} where $\\vec{n}$ is the line of sight vector and the subscript $r$ labels quantities at recombination.  \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=8.4cm, width=7.5cm]{fig1.eps} \\includegraphics[height=8cm, width=9.5cm]{plot_all_data_TT_log_sqrt.ps} \\end{tabular} \\end{center} \\caption[fig:spettri] {\\label{fig:spettri}{\\bf Left-Top:} Angular power spectrum for CMB anisotropy (TT) and for EE polarization. The latter has been amplified 20 times to make it visible in the same plot of TT. The angular scale $\\gamma$ corresponding to multipole $\\ell$ is approximately $\\gamma(^o) = 180/\\ell$. {\\bf Left-Bottom:} Filter functions of CMB telescopes with different angular resolutions. A FWHM smaller than 1$^o$ is needed to be sensitive to the \"acoustic peaks\" due to photon-baryon oscillations in the early universe. The curves are labeled with the beam FWHM. Differential instruments will not be sensitive to multipoles $\\ell < 180/\\alpha(^o)$ where $\\alpha$ is the angular separation of the beam switch; experiments scanning a limited sky region with angular size $\\theta$ cannot be sensitive to multipoles with $\\ell < 180/\\theta(^o)$ . {\\bf Right:} Selected recent measurements of the angular power spectrum of CMB anisotropy. } \\end{figure}  Our description of fluctuations with respect to the FRW isotropic and homogeneous metric is totally statistical. So we are not able to forecast the map $\\Delta T / T$ as a function of $(\\theta, \\phi)$, but we are able to predict its statistical properties. If the fluctuations are random and Gaussian, all the information encoded in the image is contained in the angular power spectrum of the map, detailing the contributions  of the different angular scales to the fluctuations in the map. In other words, the power spectrum of the image of the CMB details the relative abundance of the spots with different angular scales.  If we expand the temperature of the CMB in spherical harmonics, we have \\begin{equation} \\label{multi} {\\Delta T \\over T} = \\sum a_{\\ell,m}^T Y_\\ell^m (\\theta, \\phi) \\end{equation} the power spectrum of the CMB temperature anisotropy is defined as \\begin{equation} c_\\ell^{TT} = \\langle TT \\rangle = \\langle a_{\\ell,m}^T a_{\\ell,m}^{T,*} \\rangle \\end{equation} with no dependence on $m$ since there are no preferred directions. Since we have only a statistical description of the observable, the precision with which the theory can be compared to measurement is limited both by experimental errors {\\sl and} by the statistical uncertainty in the theory itself. Each observable has an associated cosmic and sampling variance, which depends on how many independent samples can be observed in the sky. In the case of the $c_\\ell$s, their distribution is a $\\chi^2$ with $2\\ell+1$ degrees of freedom, which means that low multipoles have a larger intrinsic variance than high multipoles (see e.g. \\cite{Whit93}).  Theory predicts the angular power spectrum of CMB anisotropy with remarkable detail, given a model for the generation of density fluctuations in the Universe, and a set of parameters describing the background cosmology. Assuming scale-invariant initial density fluctuations, the main features of the power spectrum $c_\\ell^{TT}$ are a $1/\\left[\\ell(\\ell + 1) \\right]$ trend at low multipoles, produced by the Sachs-Wolfe effect\\cite{Sach67} (second term in equation \\ref{dtt}); a sequence of peaks and dips at multipoles above $\\ell=100$, produced by acoustic fluctuations in the primeval plasma of photons and baryons\\cite{Suny70,Peeb70,Hu96}, and a damping tail at high multipoles, due to the finite depth of the recombination and free-streaming effects\\cite{Hu97a}. Detailed models and codes are available to compute the angular power spectrum of the CMB image (see e.g. \\cite{Hu02,Lewi00}). The power spectrum $c_\\ell^{TT}$ derived from the current best-fit cosmological model is plotted in fig.\\ref{fig:spettri} (top panel).  High signal-to-noise maps of the CMB have been obtained since year 2000 \\cite{debe00} (see fig.\\ref{fig:mappa}). \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=6.5cm, angle=90]{mappa.eps} \\includegraphics[height=10.5cm, angle=90]{mappawmap.eps} \\end{tabular} \\end{center} \\caption[fig:mappa] {\\label{fig:mappa} {\\bf Left:} The first map of the CMB with angular resolution and signal to noise ratio sufficient to resolve degree-sized causal horizons in the early universe was obtained by the BOOMERanG experiment, at 145 GHz, using an off-axis telescope flown on a stratospheric balloon\\cite{debe00}. The structures visible in the map are CMB anisotropies, while the contamination from local foregrounds and from instrument noise are both negligible. {\\bf Right:} The WMAP satellite has mapped the whole sky, confirming the ubiquitous presence of causal horizons, and allowing a precise determination of the power spectrum of CMB anisotropy and of the cosmological parameters \\cite{Benn03,Lars11}.} \\end{figure} From such maps, the power spectrum of CMB anisotropy is now measured quite well (see e.g. \\cite{Wrig91,debe00,Lee01,debe02,Halv02,Ruhl03,Grai03, Jone06,Kuo07,Hins03,Hins07,Nolt09,Kuo09,Fowl10,Lars11,Das11} and fig.\\ref{fig:spettri}, right panel); moreover, higher order statistics are now being measured with the accuracy required to constrain cosmological parameters (see e.g. \\cite{Das11b}). Despite of the very small signals, the measurements from independent experiments, using diverse experimental techniques, are remarkably consistent. Moreover, an adiabatic inflationary model, with cold dark matter and a cosmological constant, fits very well the measured data (see e.g. \\cite{debe94, Bond98, Bond00, Dode00, Tegm00a,Tegm00b,Brid01,Dous01,Lang01,Jaff01,Lewi02, Nett02, Ruhl03, Sper03, Benn03, Tegm04, Sper07, Koma09, Koma11}).  The CMB is expected to be slightly polarized, since most of the CMB photons undergo a last Thomson scattering at recombination, and the radiation distribution around the scattering centers is slightly anisotropic. Any quadrupole anisotropy in the incoming distribution produces linear polarization in the scattered radiation. The main term of the local anisotropy due to density (scalar) fluctuations is dipole, while the quadrupole term is much smaller.  For this reason the expected polarization is quite weak (\\cite{Rees68,Kais83,Hu97,Kami97}). The polarization field can be expanded into a curl-free component (E-modes) and a curl component (B-modes). Six auto and cross power spectra can be obtained from these components: $\\langle TT \\rangle$, $\\langle TE \\rangle$, $ \\langle EE \\rangle$, $\\langle BB \\rangle$, $\\langle TB \\rangle$, and $\\langle EB \\rangle$. For example  \\begin{equation} c_\\ell^{TE} = \\langle TE \\rangle = \\langle a_{\\ell,m}^T a_{\\ell,m}^{E,*} \\rangle \\end{equation}  \\noindent where the $a_{\\ell,m}^E$ and $a_{\\ell,m}^B$ decompose the map of the Stokes parameters $Q$ and $U$ of linear polarization in spin-2 spherical harmonics:  \\begin{equation} (Q \\pm i U)(\\theta, \\phi) = \\sum_{\\ell m} \\left(a_{\\ell m}^E \\mp i a_{\\ell m}^B \\right)  ~_{\\pm 2}Y_{\\ell m}(\\theta, \\phi) \\end{equation}.   Due to the parity properties of these components, standard cosmological models have $\\langle TB \\rangle=0$ and $\\langle EB \\rangle=0$. Linear scalar (density) perturbations can only produce E-modes of polarization (see e.g. \\cite{Selj97}). In the concordance model, $\\langle EE \\rangle \\sim 0.01 \\langle TT \\rangle$, making $\\langle EE \\rangle$ a very difficult observable to measure. The power spectrum $c_\\ell^{EE}$ derived from the current best-fit cosmological model is plotted in fig.\\ref{fig:spettri} (top-left panel).  Tensor perturbations (gravitational waves) produce both E-modes and B-modes. If inflation  happened (see e.g. \\cite{Mukh81,Guth82,Lind83,Kolb90}), it produced a weak background of gravitational waves. The resulting level of the B-modes depends on the energy scale of inflation, but is in general very weak (see e.g. \\cite{Cope93,Turn93}). Alternative scenarios, like the cyclic model\\cite{Stei02}, do not produce B-modes at all\\cite{Boyl04}.  There is a strong interest in measuring CMB polarization, and in particular the B-modes, because their detection would represent the final confirmation of the inflation hypothesis, and their level would constrain the energy-scale of the inflation process, which, we know, happened at extremely high energies (which cannot be investigated on earth laboratories \\cite{Lidd94}).  For long time attempts to measure CMB polarization resulted in upper limits (see e.g. \\cite{Cade78,Nano79,Lubi81,Masi86,Part88,Woll97}). The possibility of detecting the $\\langle BB \\rangle$ signature of the inflationary gravity waves background renewed the interest in these measurements (\\cite{Keat01,Subr00,Hedm02,Picc02,Dela02,Masi02,Vill02,Kova02,John03,Keat03a,Keat03b, Kogu03,Fare04,Cort04,Cart05}). The first statistically significant detections of CMB polarization have been reported by the coherent radiometer experiments DASI\\cite{Leit05}, CAPMAP\\cite{Bark05}, CBI\\cite{Read04}, WMAP for both $\\langle TE \\rangle$\\cite{Kogu03} and $\\langle EE \\rangle$\\cite{Page07}, and by the bolometric instrument BOOMERanG-03\\cite{Masi06,Piac06,Mont06}. The quality of CMB polarization measurements has improved steadily with the introduction of instruments with detectors arrays, like QUAD\\cite{Quad08,Quad09a,Quad09b}, BICEP\\cite{Chia10}, and QUIET\\cite{Bisc11}.  Recent measurements of the angular power spectra of CMB polarization are collected in fig.\\ref{fig:spettripol}. The polarization power spectra measured by these experiments are all consistent with the forecast from the ``concordance\" model best fitting the \\wmap $\\langle TT \\rangle$ power spectrum. In addition, they constrain the optical depth to reionization (the process ionizing the universe when the first massive stars formed), which is not well constrained by anisotropy measurements alone (see e.g.\\cite{Mort08}).   \\begin{figure} \\begin{center} \\begin{tabular}{c} \\includegraphics[height=6.5cm, width=5.5cm]{plot_all_data_TE_log.ps} \\includegraphics[height=6.5cm, width=5.5cm]{plot_all_data_EE_log_sqrt.ps} \\includegraphics[height=6.5cm, width=5.5cm]{plot_all_data_BB.ps} \\end{tabular} \\end{center} \\caption[fig:spettripol] {\\label{fig:spettripol}{\\bf Left:} Recent selected measurements of the angular cross-spectrum Temperature-E-modes-polarization $\\langle TE \\rangle$. {\\bf Center:} Recent selected measurements of the angular power spectrum of E-modes-polarization $\\langle EE \\rangle$. {\\bf Right:} Recent selected measurements (upper limits) of the angular power spectrum of B-modes-polarization $\\langle BB \\rangle$. Note the different vertical scales for the three plots. The dashed line is the model prediction for the same cosmological parameters best fitting $\\langle TT \\rangle$ measurements.} \\end{figure}  The WMAP data have sufficient coverage to allow a stacking analysis and show the irrotational pattern of polarization pseudovectors around cold and hot spots of the CMB sky\\cite{Koma11}: a clear visual demonstration of the polarization produced by density perturbations in the early universe.  To date, measurements of the rotational component of the polarization field $\\langle BB \\rangle$ resulted in upper limits, implying a ratio of tensor to scalar fluctuations $r \\simlt 0.3$. ",
        "conclusions": "\\label{sec:conclusions}  The future of CMB studies is bright. A large community has grown around the success of CMB missions, producing large amounts of excellent data. The experiments have drifted from a situation where sensitivity was the issue to a situation where control of systematic effects is the main problem. So we are facing very difficult challenges, with the ambition of understanding the most distant phenomena happening in our universe, analyzing tiny signals embedded in an overwhelming noisy background. When we approached CMB research for the first time, in 1980, measuring the intrinsic anisotropy of the CMB was considered almost science-fiction. Today CMB anisotropy is measured in a single pass with scanning telescopes using large arrays of bolometers. This experience makes us confident that much more is coming in this field, with the enthusiastic contribution of young researchers and the cross-fertilization between cosmologists, astrophysicists, solid-state / detector physicists, optics experts."
    },
    "1208/1208.0814.txt": {
        "abstract": "\\noindent GJ 581d is a potentially habitable super-Earth in the multiple system of exoplanets orbiting a nearby M dwarf. We explore this planet's long-term dynamics, with an emphasis on its probable final rotation states acquired via tidal interaction with the host. The published radial velocities for the star are re-analysed with a benchmark planet detection algorithm, to confirm that there is no evidence for the recently proposed two additional planets (f and g). Limiting the scope to the four originally detected planets, we assess the dynamical stability of the system and find bounded chaos in the orbital motion. For the planet $\\,$d$\\,$, the characteristic Lyapunov time is 38 yr. Long-term numerical integration reveals that the system of four planets is stable, with the eccentricity of the planet d changing quasi-periodically in a tight range around 0.27, and with its semimajor axis varying only a little. The spin-orbit interaction of GJ 581d with its host star is dominated by the tides exerted by the star on the planet. We model this interaction, assuming a terrestrial composition of the mantle. Besides the triaxiality-caused torque and the secular part of the tidal torque, which are conventionally included into the equation of motion, we also include the tidal torques' oscillating components. It turns out that, dependent on the mantle temperature, the planet gets trapped into the 2:1 or an even higher spin-orbit resonance. It is very improbable that the planet could have reached the 1:1 resonance. This enhances the possibility of the planet being suitable for sustained life ",
        "introduction": "\\label{firstpage} \\label{intro.sec}    Habitability of an emerging class of super-Earths -- exoplanets with masses greater than that of the Earth but lower than that of Uranus -- depends on a combination of physical parameters. Among these, crucial is the intensity of irradiation from the host star. A favourable rate of irradiation permits water to be available in the liquid form. The irradiation intensity depends on the luminosity of the star, the size of the planet's orbit and, to a lesser degree, on the orbital eccentricity.  The chemical composition of the planet's atmosphere influences the average temperature and the temperature variations on the surface. For example, estimates show that a certain minimal amount of CO$_2$ in the atmosphere of GJ 581d is required to keep water above the freezing point on the surface, while GJ 581c is likely to have experienced a runaway greenhouse event \\citep{vonP,hu}. Planet GJ 581d, which is the main target of our study, may be located on the outer edge of the habitable zone, according to \\citet{vonB}.  We begin our study with addressing the still controversial problem of composition of this planetary system. In Section \\ref{orb.sec}, we confirm that only the four originally detected planets (b through e) are real and detectable at the 0.99 confidence level by our fully automated, fast grid search algorithm. The two additional planets, f and g, proposed by \\citet{vog} are not found with any combination of the published radial velocity (RV) data.  The orbit estimation technique employed in our paper is designed, in particular, to produce accurate and robust results for the eccentricity of each detected planet -- an important asset for the subsequent dynamical simulations. The dynamical stability and the presence of chaos in the orbital motion of the four planets is investigated in Section \\ref{cha.sec} by long-term integrations, assuming zero inclination and coplanar orbits. The system is found to be long-term stable but strongly chaotic, with the eccentricity and semimajor axis of planet d little varying over gigayears. This allows us to investigate, in Section \\ref{tid.sec}, the spin-orbit dynamics of GJ 581d with its orbital parameters fixed. The planet is assumed to have a terrestrial rheology and a near-zero obliquity. ",
        "conclusions": "The growing number of detected systems of multiple exoplanets and the impressive quality of observational data collected for them allow astronomers to perform analysis of probable dynamical states and evolution of these remote worlds at a level of detail unthinkable just several years ago. Still, this analysis is riddled with difficulties and uncertainties. The planets of GJ 581 present a challenge for both observational practice and interpretational theory. The story of two fictitious planets tells a lesson about the hazards of combining, without proper caution, the data from two instruments with their own sets of systematic errors. It also calls for a certain unification of the planet detection techniques or, at least, for an easily accessible and well-tested standard detection algorithm available as a web application. It is fine to apply a variety of different detection methods to the same data, but the standard algorithm should always be checked, and discrepancies, if any, should be investigated and reported. $N$-body integration of detected systems should also be a norm reducing the probability of error. The system of GJ 581 proves to be remarkably stable, with the four {\\it{bona fide}} planets remaining on their orbits despite the strong evidence of chaos. The characteristic Lyapunov times are very short compared to the dynamical lifetime of the system.  We have explored the rotation history of the planet GJ 581d assumed to have composition alike to that of the terrestrial planets of the Solar system.  Contrary to the previous publications on the subject, which were based on {\\it{ad hoc}}, simplistic tidal theories, we find that this planet {\\underline{cannot}} be captured into a synchronous or pseudo-synchronous rotation, if it began its evolution from faster, prograde initial spin rates. Instead, for a plausible range of parameters, the most likely state of the planet is the 2:1 spin-orbit resonance. In this state, the day on GJ 581d should be 67 Earth's days long, which bides for an inhospitable environment, though the potential habitability of this planet cannot be ruled out just on climatic considerations. The case of 2:1 spin-orbit resonance was considered in the simulations of a hypothetical atmosphere on GJ 581d by \\citet{Wordsworth} whose modeling confirmed the possibility of liquid water being present on the surface, under some favorite conditions.  The next likeliest equilibrium states are the 3:2 or 5:2 resonances, depending on the temperature and viscosity of the mantle (much less on the planet's triaxiality).  At the same time, in the event that the initial rotation of the planet was retrograde, the most probable final state is synchronous rotation."
    },
    "1208/1208.0559_arXiv.txt": {
        "abstract": "We analyse the situation where the primordial curvature perturbations are produced by the joint effects of an inflaton field and two curvaton fields.  We present general equations which allow the reader to obtain $f_{NL}$ for several different scenarios which differ in the order in which fields decay into radiation after inflation.  In order to investigate the physics of these equations we analyse some simplified situations where the fields are harmonic and both curvatons are frozen at the same expectation value during inflation.  We find quite complex behaviour -  for a given situation where the inflaton contributes a fixed amount to the total curvature perturbation there are situations where $f_{NL}$ is maximised if both curvatons share equally in contributing the rest and situations where $f_{NL}$ is maximised if only one of the two curvatons contributes the rest.  We are unable therefore to make any completely general extrapolations about the expected non-gaussianity from $N$ curvatons.  We find that as the curvaton contribution to the overall perturbation is gradually increased, $f_{NL}$ rises to a maximum before falling again and that a given $f_{NL}$ can correspond to many different parameter sets for the two curvatons. ",
        "introduction": "Inflation is the most successful theory we have to describe the early universe. According to inflation \\cite{liddle}, at some early time the universe underwent an accelerated expansion, and this solves the horizon and flatness problem while also predicting an almost scale-invariant power spectrum in the cosmic microwave background, consistent with current observations. The most basic models of inflation (see, for example, \\cite{lazarides}) involve a single, massive scalar field $\\phi$ called the inflaton, which slowly rolls down its potential to source the necessary expansion. Once it reaches the minimum of the potential it oscillates around that point, reheating the Universe by decaying into standard model particles. The quantum fluctuations of the inflaton get stretched out with the expansion and become classical perturbations, causing inhomogeneities which upon horizon re-entry evolve to become the large-scale structure that we observe today \\cite{liddlelyth}.  Since many theories that go beyond the standard model predict scalar fields, it is not unreasonable to assume that in early times, these fields had expectation values away from the minimum of their potential. Such fields, known as curvaton fields, can be present during inflation but only begin to move when the Hubble parameter drops down to their mass (as in \\cite{lythwands} and \\cite{lythungarelliwands}). They then proceed to oscillate around the minimum of their potential and eventually decay when the Hubble parameter is equal to their decay rate. The presence of these particles imprints a record of the curvaton perturbations as well as the inflaton perturbations upon the cosmic distribution of matter and therefore provides a different mechanism for the creation of the primordial curvature perturbation $\\zeta$ and can alter the predictions of basic single-field inflation \\cite{bartololiddle}, \\cite{wandsmaliklythliddle}, \\cite{byrneschoi}, \\cite{wands}, \\cite{dimopoulos}. There are a plethora of curvaton models one can imagine existing in particle physics and string theory and they typically include a large number of parameters. To constrain them, one needs to compare their predictions to current astrophysical observations, with the most popular discriminant being the non-gaussianity parameter, $f_{NL}$ (\\cite{langlois}, \\cite{bartolo},\\cite{gordonlewis}, \\cite{komatsu}).  The idea behind the curvaton scenario is to boost small perturbations and make them larger by virtue of the energy density in the curvaton field red-shifting more slowly than the energy density from the inflaton field.  This occurs once the inflaton has decayed into radiation but the curvaton is still oscillating (and thus red-shifting like matter).  There are therefore multiple scenarios depending upon the order in which the inflaton decays, the curvaton starts to move and the curvaton decays.  In this work we will attempt to study a number of these different situations, although we will not consider the case where the curvaton starts to roll during inflation, as this moves into the territory of multifield inflation.  In inflation models, the non-gaussianity parameter $f_{NL}$ is related to the curvature perturbation after the last decay. This parameter is a measure of deviations from a purely Gaussian distribution of the primordial curvature perturbation and the current $f_{NL}$ limits are given by WMAP 7 to be $-10<f_{NL}<74$ \\cite{wmap7}. We analytically compute the expression for $\\zeta$ at the last decay and $f_{NL}$ in models with one inflaton, $\\phi$, and two curvatons, $\\sigma$ and $\\chi$, for various sequences of particle decays. In each case, the inflaton is the one responsible for producing the necessary e-folds and both the curvatons are frozen during the inflaton's slow-roll. Once inflation has finished, $\\phi$  begins to oscillate around the minimum of its potential and then decays. Independently, some time after the end of inflation, when $H = m_{curvaton}$, the curvatons begin to oscillate and eventually decay.  All three particles contribute to the curvature perturbation.  We will evaluate the primordial curvature perturbation at the time of the last particle's decay, whether it is the inflaton or one of the curvatons. ",
        "conclusions": "In this work, we have studied the non-gaussianity predicted in models of inflation with 1 inflaton and 2 curvatons. We derived a general expression for $f_{NL}$ from the bi-spectrum of a general curvature perturbation $\\zeta$. We saw that our general expression simplifies to reproduce other expressions for $f_{NL}$ in the literature under the relevant simplifications.  We then analytically derived the necessary coefficients for specific scenarios in which we varied the sequence of the decays of the three particle species. For our main three cases, we then simplified our $f_{NL}$ expression under the assumption that the last two decays happen at the same time.  For all of the calculations we performed in this paper, we assumed a harmonic potential for the inflaton of the form $m^2\\phi^2$ and we assumed therefore a slow roll parameter 60 efolds before the end of inflation of $\\epsilon\\sim 0.02$. We made no assumption for the potential of the curvatons in our general calculations; for the simplified cases we used the quadratic form.  In the simultaneous decay limit we studied the features of the predicted $f_{NL}$ and found that the produced non-gaussianity is bounded. When the contribution of the curvaton to the final perturbation is of order unity, $f_{NL}$ is slightly negative but becomes more positive as $f_{\\sigma_\\alpha}$ and $f_{\\chi_\\beta}$ become smaller.  The maximum value of $f_{NL}$ is set in this situation by the expectation value $\\sigma_*$ of the curvaton during inflation and is smaller when $\\sigma_*$ is larger.  For expectation values smaller than around $10^{16}$ GeV $f_{NL}$ grows beyond the current observational constraints if the relative contribution of the curvaton has the right value. If the expectation value of the curvaton is $10^{17}$ GeV then we find that the maximum non-gaussianity is around $f_{NL}\\sim 0.8$ which is not in the range that will be detectable by Planck.  We find that $f_{NL}$ can take very large values, depending upon the expectation value of the curvaton during inflation but that it does not grow indefinitely. After it reaches a maximum it decreases and returns to within the range allowed by current observations.  A particular value of $f_{NL}$ therefore corresponds to two different sets of parameters, even in the situation where there is a single inflaton and a single curvaton and more information is required.  The predicted $f_{NL}$ in all of the three field models in this paper is quite a complicated function of the relative contributions of the two curvatons.  If, for example, we take the situation where 80\\% of the total curvature perturbation comes from the inflaton and 20\\% comes from the combined effect of the two curvatons, there are situations where having all of the curvaton contribution to the pertubation coming from one of the two curvaton fields will maximise $f_{NL}$, and there are situations where having 10\\% of the perturbation coming from each curvaton will maximise $f_{NL}$, depending upon the expectation values of the curvaton fields.  There are therefore in general several combinations of field parameters that give the same non-gaussianity and, alternatively, even when these parameters change in a way that leaves the total curvaton contribution to $\\zeta$ constant there can be big changes in the value of $f_{NL}$.  Of course it would be interesting therefore to investigate the trispectrum to see if the parameter $g_{NL}$ could contribute to distinguishing between different models.  It is also quite clear that the machinery used in this paper leads to rather a lot of long equations.  It would be nice to come up with a way of simplifying the calculations such that one could make more general statements about N curvatons."
    },
    "1208/1208.0135_arXiv.txt": {
        "abstract": "We have obtained time series observations of the Orion Nebula Cluster at 70$\\,\\mu$m and 160$\\,\\mu$m from the \\emph{Herschel}/PACS Photometer. This represents the first wide-field far-infrared photometric monitoring of a young star forming region. The acquired 35\\arcmin$\\times$35\\arcmin~maps show complex extended structures, with unprecedented details, that trace the interaction between the molecular gas and the young hot stars. We detect 43~protostars, most of which are situated along the integral-shaped filament extending from the Orion nebula, through OMC$\\,$2 and to OMC$\\,$3. We present high-reliability light curves for some of these objects using the first six epochs of our observing program spread over 6~weeks. We find amplitude variations in excess of 20\\% for a fraction of the detected protostars over periods as short as a few weeks. This is inconsistent with the dynamical time-scales of cool far-IR emitting material that orbits at hundreds of AU from the protostar, and it suggests that the mechanism(s) responsible for the observed variability originates from the inner region of the protostars, likely driven by variable mass accretion. ",
        "introduction": "\\label{sec:intro}  Photometric variability was one of the original, defining characteristics of the class of objects that were later determined to be stars in the process of formation, or Young Stellar Objects (YSOs) \\citep{joy}. While optical and near-IR variability studies have demonstrated the relation between accretion shocks and hotspots on the surface of rotating YSOs \\citep[e.g.,][]{vrba, carpenter}, observations at longer wavelengths offer another perspective to study cooler circumstellar disks and envelopes surrounding nascent stars. In particular, recent mid-IR photometric and spectroscopic time series of YSOs with the \\emph{Spitzer} Space Telescope \\citep{werner} have begun to shed light on inner disk structures, identifying both inner disk warps and `clouds' in the disk \\citep[e.g.,][]{morales09, muzerolle, espaillat}. Elaborate models offer satisfying explanations to the observed time series \\citep{flaherty, turner, flaherty11, ke}.  Similar photometric variability studies are a lot sparser in the literature at longer wavelengths ($\\lambda > 50\\,\\mu$m), mostly due to the difficult access to this wavelength domain. \\citet{harvey} report on a 200\\% flux variation of SSV$\\,$13 in NGC$\\,$1333 over a two-year period using the Kuiper Airborne Observatory. \\citet{juhasz} and \\citet{sitko} analyze the 1-300$\\,\\mu$m variability of the star SV$\\,$Cep over two years, and they invoke a growing warp at the inner edge of the disk, which leads to the shadowing of the outer disk, to reproduce the observations across the entire wavelength range. Based on reprocessed ISOPHOT observations \\citep{lemke}, \\citet{kospal} present the far-IR fading of the pre-main sequence star OO$\\,$Serpentis after it went into outburst, and they show that the $60\\,\\mu$m flux changed by a factor of 8 over a period of 11~years. Most sources targeted by these monitoring programs are disk-bearing YSOs.  In this article, we present the first sensitive wide-field far-IR photometric monitoring of protostars which significantly extends the statistics on YSO variability monitoring in this wavelength regime and evolutionary stage. We recently completed the first set of time series observations of the Orion Nebula Cluster (ONC) with the PACS instrument \\citep{poglitsch} onboard the \\emph{Herschel} Space Observatory \\citep{pilbratt}. During the first two-month visibility window, we obtained six observations of a single field centered on the integral-shaped filament located north of the ONC, which contains hundreds of YSOs. In addition, this field has been monitored in the optical \\citep{herbst}, and in the mid-IR with \\emph{Spitzer} as part of the YSOVAR program \\citep{morales11}, thus making it the largest YSO variability database to date.  We first describe the observations and data reduction in section~\\ref{sec:method}, then we give a detailed account of the photometric measurements and present the light-curves extracted from the time series observations. Finally we discuss the possible origin of the observed variability in section~\\ref{sec:result}. ",
        "conclusions": "\\label{sec:result}  \\begin{figure*} \\centering \\begin{tabular}{cc} \\includegraphics[width=0.48\\textwidth]{lightCurve_blue_ID10_dual_lineaCorrDebugRec.eps}  & \\includegraphics[width=0.48\\textwidth]{lightCurve_blue_ID25_dual_lineaCorrDebugRec.eps} \\\\ \\includegraphics[width=0.48\\textwidth]{lightCurve_blue_ID22_lineaCorrDebugRec.eps}  & \\includegraphics[width=0.48\\textwidth]{lightCurve_blue_ID6_lineaCorrDebugRec.eps} \\\\ \\includegraphics[width=0.48\\textwidth]{lightCurve_blue_ID3_lineaCorrDebugRec.eps} & \\includegraphics[width=0.48\\textwidth]{lightCurve_blue_ID23_lineaCorrDebugRec.eps} \\\\ \\end{tabular} \\caption{Sample of reliable PACS light curves drawn from Table~\\ref{tab:fluxes_LC}. The left column presents the light curves of variable protostars, and the right column contains those that show flux variations within the estimated photometric uncertainties (see section~\\ref{sec:photometry} for details). The set of graphs in the top row show the light curves at 70\\,$\\mu$m and 160\\,$\\mu$m when both are deemed reliable, while the other plots give 70\\,$\\mu$m fluxes only. The horizontal dashed and dotted lines give the average flux of the sources and the $\\pm$5\\% variations around the mean, respectively, indicating our level of confidence for variability detections. The first epoch was obtained on Feb. 26, 2011, corresponding to the Herschel Operational Day 653, or MJD~55618.} \\label{fig:lightCurves} \\end{figure*} \\begin{figure*} \\centering \\begin{tabular}{cc} \\includegraphics[width=0.48\\textwidth]{lightCurve_Spitzer_1433.eps}  & \\includegraphics[width=0.48\\textwidth]{lightCurve_Spitzer_1321.eps} \\\\ \\includegraphics[width=0.48\\textwidth]{lightCurve_Spitzer_1339.eps}  & \\includegraphics[width=0.48\\textwidth]{lightCurve_Spitzer_1449.eps} \\\\ \\includegraphics[width=0.48\\textwidth]{lightCurve_Spitzer_1456.eps} & \\includegraphics[width=0.48\\textwidth]{lightCurve_Spitzer_1300.eps} \\\\ \\end{tabular} \\caption{Spitzer/IRAC light curves measured at 4.5\\,$\\mu$m on the 6~objects presented in figure~\\ref{fig:lightCurves}. These observations were obtained as part of the YSOVAR program in 2009, they are not contemporaneous with the Herschel data. The plots are arranged as in figure~\\ref{fig:lightCurves} with the YSOVAR name given in the plot title, and the horizontal dashed and dotted lines also represent the mean and $\\pm$5\\% variations, respectively. } \\label{fig:YSOVARlightCurves} \\end{figure*}  It is a difficult task to disentangle intrinsic source variability and photometric errors based solely on 6~data points, especially for faint sources. Nevertheless, considering the criterions we used in selecting reliable light curves (cf section~\\ref{sec:photometry}), we expect that flux variations greater than $\\sim$10\\% can be confidently attributed to the intrinsic variability of the observed sources.%  We find that 8~sources out of 17~show peak-to-peak flux variations higher than $\\sim$10\\%. The left column of figure~\\ref{fig:lightCurves} presents a sample of such light curves. The case of HOY$\\,$J053522.27-050116.8 (top left plot) is of particular interest with its sine-like light curve that shows an amplitude of nearly 20\\%. The smooth shape of the light curve suggests that the photometric uncertainties are very small for this source. In addition both 70$\\,\\mu$m and 160$\\,\\mu$m light curves show the same trend, which strengthens further our confidence that this source is indeed variable in the Far-IR. The object HOY$\\,$J053524.23-050831.9 also shows strong variations in excess of 20\\% in just over 2~weeks. For comparison, we present in figure~\\ref{fig:YSOVARlightCurves} Spitzer light curves obtained in 2009 as part of the YSOVAR program \\citep{morales11}, and it appears that objects variable in the far-IR tend to be also variable in the mid-IR (left column of figures~\\ref{fig:lightCurves} and~\\ref{fig:YSOVARlightCurves}). However the shape of Herschel and Spitzer light curves are not directly comparable since they are not contemporaneous\\footnote{There was no significant overlap in the Spitzer and Herschel visibility windows of Orion in 2011.} and they have significantly different observing cadences.  The remaining 9~sources that have reliable light curves show flux variations in the 5 -- 10\\% range. Although those sources exhibit variations below the 10\\% threshold, we cannot exclude that they might as well be intrinsically variable in the far-IR. In fact, based on the Herschel dataset currently available, we are only sensitive to variations with timescales in the range 10 -- 50~days, and we miss all variations on shorter and longer timescales. For instance, the light curve of HOY~J053524.71-051030.6 does not satisfy our variability criterion according to its 70$\\,\\mu$m flux only, however its Spitzer light curve shows clear variability in the mid-IR with a period of $\\sim$4~days (bottom right plot of figure~\\ref{fig:YSOVARlightCurves}). If the mid- and far-IR variability had the same origin, and thus showed similar timescales at both wavelengths, then the sampling frequency of 10~days would not be appropriate to detect such rapid variations in Herschel bands. \\\\  Although YSOs have been monitored on daily to yearly timescales from the optical to the mid-IR and in all evolutionary stages \\citep[e.g. ][]{herbst, morales11}, the variability parameter-space is still sparsely sampled in the far-IR. This is particularly true for high cadence monitoring and embedded protostars, as most far-IR monitoring programs have targeted outbursting YSOs with fading timescales of several years \\citep[e.g. ][]{kospal}. The present Herschel time-series observations thus fill in this parameter-space gap, and they show that the far-IR emission, which is a good tracer of the internal luminosity of protostars \\citep{dunham}, can vary noticeably on timescales as short as a couple of weeks. This is orders of magnitude shorter than the dynamical timescales of the far-IR emitting material that orbits at hundreds to thousands of AUs from the central protostar. Such short timescales indicate that the mechanism(s) responsible for the observed variability occurs on smaller spatial scales, either at the surface or within the disk, or close to the central protostar (r $\\ll1\\,$AU) where timescales are consistent with weekly to monthly timescales.  Non-steady accretion likely plays a role in the observed variability. \\citet[][and references therein]{zhu} argue that the typical accretion luminosity observed in protostars implies an accretion rate considerably lower than the predicted time-average infall rate in YSOs. This `luminosity problem' can be explained if infalling material accretes sporadically causing major accretion events that are sufficiently short-lived so that protostars are usually observed in quiescence. Numerical models of \\citet{vorobyov} manage to reproduce these accretion events for Class\\,I and Class\\,II YSOs when disk fragments fall onto the central protostar. Such energetic outbursts have been observed in FU-Orionis-type objects with typical timescales of 1 -- 10~years \\citep{hartmann}. Similar events might occur on shorter temporal and less energetic scales.  For instance, \\citet{flaherty11} measured accretion rates from Pa$\\beta$ and Br$\\gamma$ observations, and they find that the accretion rate of the Class$\\,$II object~LRLL~31 varies by a factor of five, with the largest changes occurring on a weekly timescale. These episodic accretion events could be pictured as knotty filaments of gas funneling from the inner region of the circumstellar disk to the central star, likely through magnetic field lines \\citep{shu}, and releasing packets of energy as the clumps hit the surface of the protostar. For embedded protostars, the energetic photons generated in the accretion shock are reprocessed by the dense and dusty surrounding envelop, and they eventually escape the protostar in the far-IR regime. The accretion luminosity variations that originate from episodic accretion events in the central region of a protostar propagate outwards and temporarily warm up the envelop, which in turn leads to detectable far-IR flux variations. The typical timescale for photons to reach the far-IR emitting material in the envelop (located at $\\sim10^2-10^3\\,$AU from the central star) is roughly the light travel time, which is a few days, similar to the typical timescales observed in the light curves of Figure~\\ref{fig:lightCurves}.   The alternative scenario proposed by \\citet{flaherty} might also cause a measurable far-IR variability in YSOs. It consists in a variable scale height of the disk inner edge that casts a shadow on the outer disk, thus cooling down the mid- and far-IR-emitting disk. This non-axisymetric disk model was used to reproduce the mid-IR variability observed in the YSO LRLL~31, and it can potentially give constraints on the inner disk structure. In a subsequent work, \\citet{flaherty11} have tested various models that may lead to variable scale heights of the inner edge of the disk, with origins ranging from variable accretion, perturbations by a companion, winds, and the influence of magnetic fields. Although LRLL~31 is more evolved than most objects observed in the present study, a similar mechanism could well explain the variations of the Class\\,II source HOY\\,J053530.33--045938.6 (bottom left plot in figure~\\ref{fig:lightCurves}). However, contemporaneous near-IR observations would be necessary to confirm the expected anti-correlated flux variations with respect to the far-IR light curves.\\\\  The first 6~epochs of our monitoring program have demonstrated that protostellar emission can vary on relatively short timescales in the far-IR. Furthermore we find that the fraction of variable protostars in our sample is relatively high ($>$40\\%), though difficult to ascertain based on sparse 70$\\,\\mu$m light curves only. For a better sampling of the light curves, we have requested a higher observing cadence for the second Orion visibility window, and we have extended the initial span of the monitoring program with the last observations scheduled for autumn~2012, thus covering a period of nearly 2~years. These additional observations, with refined photometric measurements and a detailed SED modeling, should increase the number of variability detections and help us understand the nature of these objects and the mechanism responsible for the observed variability. A spectroscopic follow-up with Herschel/PACS is also scheduled for 2012 to monitor accretion activity in a couple of variable sources with atomic and molecular lines, particularly [OI], [CII], CO and H$_2$O, already observed as part of the HOPS program \\citep{megeath,manoj}.\\\\    This work is based on observations made with Herschel, a European Space Agency Cornerstone Mission with significant participation by NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech. This research has made use of NASA's Astrophysics Data System.  {\\it Facilities:} \\facility{Herschel Space Observatory (PACS)}."
    },
    "1208/1208.1419_arXiv.txt": {
        "abstract": "We report on an observation of SGR~1627$-$41 made with the {\\em Chandra X-ray Observatory} on 2011 June 16. Approximately three years after its outburst activity in 2008, the source's flux has been declining, as it approaches its quiescent state. For an assumed power-law spectrum, we find that the absorbed 2--10 keV flux for the source is $1.0^{+0.3}_{-0.2} \\times 10^{-13}$ erg cm\\textsuperscript{$-$2} s\\textsuperscript{$-$1} with a photon index of $2.9 \\pm 0.8$ ($N_H=1.0\\times10^{23}$ cm\\textsuperscript{$-$2}). This flux is approximately consistent with that measured at the same time after the source's outburst in 1998. With measurements spanning 3 years after the 2008 outburst, we analyze the long-term flux and spectral evolution of the source. The flux evolution is well described by a double exponential with decay times of 0.5 $\\pm$ 0.1 and 59 $\\pm$ 6 days, and a thermal cooling model fit suggests that SGR~1627$-$41 may have a hot core ($T_c\\sim2\\times 10^8\\ {\\rm K}$). We find no clear correlation between flux and spectral hardness as found in other magnetars. We consider the quiescent X-ray luminosities of magnetars and the subset of rotation-powered pulsars with high magnetic fields ($B\\gtrsim10^{13}\\ {\\rm G}$) in relation to their spin-inferred surface magnetic-field strength, and find a possible trend between the two quantities. ",
        "introduction": "Magnetars are neutron stars with ultra-strong magnetic fields, the decay of which is theorized to power the observed radiation from the star  \\citep{td95, td96, tlk02}. Soft Gamma Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs) are two observational manifestations of magnetars. The former show repeated soft gamma-ray bursting activity and have relatively hard spectra, and the latter have soft spectra typically characterized by a blackbody plus power law. However, the distinction between the two has become increasingly blurred \\citep{tlk02, gkw02, plm+07, pp11}. Magnetars generically show X-ray outbursts, which are sudden increases of luminosity by orders of magnitude for days to months. During outbursts, almost all the properties of magnetars, such as their flux, spectrum and pulsed flux, change \\citep[see][for reviews]{wt06, k07, m08, re11}. The X-ray luminosities of the first-discovered magnetars were typically $\\sim10^{35}$ erg s\\textsuperscript{$-$1}, orders of magnitude greater than their spin-down luminosity. However, more recent discoveries of magnetars in outburst suggest most magnetars in quiescence may be far less luminous \\citep[e.g. 1E~1547.0$-$5408, XTE~J1810$-$197, Swift~J1822$-$1606; ][]{gg07, bid+09, snl+12}. The spin periods of magnetars\\footnote{See the online magnetar catalog for a compilation of known magnetar properties, http://www.physics.mcgill.ca/$\\sim$pulsar/magnetar/main.html} are in the narrow range of 2--12 s and are relatively long compared to those of radio pulsars. The magnetic-field strength inferred from the spin period and spin-down rate is typically $B > 10^{14}$ G, assuming the standard vacuum dipole formula for magnetic braking, although several with $B \\lesssim 10^{14}\\ {\\rm G}$ have recently been found \\citep{ret+10,snl+12,rie+12}.  The recent discoveries of a magnetar-like outburst from a high-$B$ rotation-powered pulsar (RPP) \\citep[PSR~J1846$-$0258,][]{ggg+08}, pulsed radio emission from a magnetar \\citep[XTE~J1810$-$197,][]{crh+06}, a low magnetic field magnetar \\citep[SGR~0418$-$5729,][]{ret+10}, and the very low X-ray luminosities measured for several magnetars are of particular interest. These raise important questions about the relationship between magnetars and other types of neutron stars, and why there is such an apparent diversity in magnetar properties. For example, what determines a magnetar's quiescent X-ray luminosity? According to conventional magnetar models \\citep{td95, td96, tlk02}, there should be a correlation between $B$ and the X-ray luminosity. However, transient magnetars, with their faint quiescent luminosities compared with non-transient magnetars of the same inferred $B$, challenge this. And relatedly, are high-$B$ RPPs in general magnetars in quiescence? PSR~J1846$-$0258's outburst suggests this, but it is only one source. \\citet{plm+07} and \\citet{pp11} argue that magnetars and RPPs are all related on the basis of magneto-thermal evolution theory. Also a possible connection between high-$B$ RPPs and magnetars has been suggested on the basis of possibly high thermal temperatures of high-$B$ RPP X-ray emission \\citep{km05, zkg+09, okl+10, zkm+11, nk11}. However, it is important to study many magnetars and RPPs in quiescence to better address this question.  SGR~1627$-$41 was discovered on 1998 June 15 with the Burst and Transient Source Experiment \\citep[BATSE,][]{Fishman1989}. It was then identified as a SGR by \\citet{kkw+98}. It is located at R.A. = 16\\textsuperscript{h}35\\textsuperscript{m}51\\textsuperscript{s}.844, Dec. = $-$47$^\\circ$35$'$23$''$.31 (J2000.0) and is estimated to be $11.0 \\pm 0.3$ kpc away based on an apparent association with a star-forming region and molecular cloud \\citep{hkw+99, ccd+99,wpk+04}. The spin period and the spin-down rate were not measured until recently due to the faint nature of the source in quiescence. After another outburst in 2008 May, the spin period and the spin-down rate were measured to be 2.594578(6) s and $1.9(4) \\times 10^{-11}$ s s$^{-1}$, which imply an inferred surface dipolar magnetic-field strength of $B \\equiv 3.2 \\times 10^{19}(P \\dot P)^{1/2}$ G $= 2 \\times 10^{14}$ G \\citep{etm+09, ebp+09}. The lowest flux ever measured for this magnetar is $6 \\times 10^{-14}$ erg cm\\textsuperscript{$-$2} s\\textsuperscript{$-$1} \\citep[2--10 keV, $\\sim$10 years after the 1998 outburst, ][]{eiz+08}. However, whether it was in quiescence at that time is not clear; the luminosity might have declined further had the outburst not occurred in 2008.  Here, we report the results of an observation made with {\\em Chandra} in the direction of SGR~1627$-$41. We then combine our flux measurement of SGR~1627$-$41 with previous values to determine the long-term flux evolutions after its outbursts, and attempt to fit them to a model to infer thermal properties of the source. We compare the flux we measure with that from the same time after the first outburst and the lowest flux ever measured. Finally, we compile quiescent X-ray luminosities of magnetars and high-$B$ RPPs to search for a correlation with spin-inferred $B$, as might be expected in the magnetar model. ",
        "conclusions": "Using {\\em Chandra} observations, we have measured the spectrum and absorbed flux in the 2--10 keV band for SGR~1627$-$41 approximately 3 years after its 2008 outburst. The spectrum was consistent with a power law having $\\Gamma = 2.9 \\pm 0.8$ (or a blackbody having $kT = 0.85^{+0.25}_{-0.16}$ keV), and the absorbed flux was $1.0^{+0.3}_{-0.2} \\times 10^{-13}$ erg cm\\textsuperscript{$-$2} s\\textsuperscript{$-$1} ($0.75^{+0.20}_{-0.17} \\times 10^{-13}$ erg cm\\textsuperscript{$-$2} s\\textsuperscript{$-$1} for a blackbody spectrum) in 2011 June. Although the source flux is similar to that detected a comparable amount of time following its 1998 outburst and is similar to the lowest yet seen from this source, it is unclear whether it has reached true quiescence, as its spectrum is significantly harder than in other magnetars in quiescence. We showed that the flux evolution of the source after its outburst activity in 2008 followed a double exponential with decay times of $0.5 \\pm 0.1$ and $59 \\pm 6$ days. Our model fitting, assuming the flux relaxation is due to crustal cooling, suggests that the core temperature of SGR~1627$-$41 is high ($T_c\\sim2\\times 10^8\\ {\\rm K}$) and that the energy was deposited in the outer crust (at different depths) for the two outbursts. This is the same conclusion as for Swift~J1822$-$1606 \\citep{snl+12} and may provide an interesting constraint on crust breaking models. We show that the 2008 activity of SGR~1627$-$41 was likely to have been initiated by a crustal fracture, causing a twist of the external magnetic fields. However, for this magnetar, we see no clear correlation between flux and spectral hardness as seen in other magnetars, which is puzzling. Finally, we find a possible correlation between the inferred magnetic field and the quiescent luminosity of 16 magnetars (including two candidates). We also note that the correlation becomes stronger if we include high-$B$ RPPs, which further suggests a connection between high-$B$ RPPs and magnetars. The discovery and detailed study of more high-$B$ RPPs and magnetars in the future will help us to better understand the physical connection between these two populations.  VMK acknowledges support from a Killam Fellowship, an NSERC Discovery Grant, the FQRNT Centre de Recherche Astrophysique du Qu\\'ebec, an R. Howard Webster Foundation Fellowship from the Canadian Institute for Advanced Research (CIFAR), the Canada Research Chairs Program and the Lorne Trottier Chair in Astrophysics and Cosmology. JAT, AB, EG, and FR acknowledge partial support from NASA through {\\it Chandra} Award Number GO1-12068A issued by the {\\it Chandra} X-ray Observatory Center, which is operated by the Smithsonian Astrophysical Observatory under NASA contract NAS8-03060. AC is supported by an NSERC Discovery Grant and the Canadian Institute for Advanced Research (CIFAR)."
    },
    "1208/1208.3907_arXiv.txt": {
        "abstract": "Due to orbital decay by gravitational-wave radiation, some close-binary helium white dwarfs are expected to merge within a Hubble time. The immediate merger products are believed to be helium-rich sdO stars, essentially helium main-sequence stars. We present new evolution calculations for these post-merger stars beyond the core helium-burning phase. The most massive He-sdO's develop a strong helium-burning shell and evolve to become helium-rich giants. We include nucleosynthesis calculations following the merger of $0.4 \\rm M_{\\odot}$ helium white-dwarf pairs with metallicities $Z = 0.0001, 0.004, 0.008$ and 0.02. The surface chemistries of the resulting giants are in partial agreement with the observed abundances of R Coronae Borealis and extreme helium stars. Such stars might represent a third, albeit rare, evolution channel for the latter, in addition to the CO+He white dwarf merger and the very-late thermal pulse channels proposed previously. We confirm a recent suggestion that lithium seen in R\\,CrB stars could form naturally during the hot phase of a merger in the presence of \\iso{3}{He} from the donor white dwarf. ",
        "introduction": "The existence of close binary white dwarfs has been predicted theoretically \\citep{Han98, Nelemans00, Nelemans01} and demonstrated observationally \\citep{Rebassa11, Brown11, Kilic10, Kilic11}. It has been demonstrated that the orbits of close binaries decay through the emission of gravitational wave radiation \\citep{Evans87, Cropper98}. Models for the merger of nearly-equal mass white dwarfs following spiral-in demonstrate that these occur on a dynamical timescale (minutes) \\citep{Benz90, Guerrero04, Pakmor11}. The less massive WD is disrupted and part of the debris forms a prompt hot corona, the remainder forms a disk which subsequently accretes onto the surviving white dwarf \\citep{Yoon07,Loren09}. Models for the evolution of the product of a double helium-white-dwarf merger demonstrate the off-centre ignition of helium burning, followed by expansion and evolution onto the helium main sequence \\citep{Saio00}. Models which include the composite nature of the debris (corona + disk) successfully account for the distribution in effective temperature, luminosity and surface composition of compact helium-rich subdwarf O stars \\citep{Zhang12}. In particular, the most-massive mergers become the hottest sdO stars with carbon-rich surfaces, whilst the least-massive mergers are cooler with predominantly nitrogen-rich surfaces.  Following inward migration of the helium-burning shell, the merged star is essentially a helium main-sequence star having a mass between 0.4 and 1.0 M$_{\\odot}$, although its surface layers will carry a record of the merger event. The evolution of low-mass helium stars has been studied for over 40 years \\citep{Paczynski71, Dinger72, Trimble73, Weiss87}. The lowest-mass stars will evolve directly to become CO/He white dwarfs following core helium exhaustion. For $M \\ge1.0 {\\rm M}_{\\odot}$, shell helium burning will ignite around the carbon/oxygen core, and the star will expand to become a giant. Early authors proposed this as a possible origin for the hydrogen-deficient R Coronae Borealis (R CrB) variables \\citep{Paczynski71, Weiss87}.  R CrB stars and the hotter extreme helium stars are low-mass supergiants of spectral types F, A and B. Their surfaces are extremely deficient in hydrogen, and enriched in carbon, oxygen, neon and other nuclear waste. Two principle evolution channels have been established. A small fraction may be produced following a late thermal pulse in a post-asympotic giant-branch star on the white dwarf cooling sequence \\citep{Iben84a, Clayton11}. The majority are more likely to have been produced following the merger of a carbon-oxygen white dwarf with a helium white dwarf \\citep{Webbink84, Saio02, Jeffery11, Longland11}.  In this letter we demonstrate that high-mass double helium-white-dwarf mergers could contribute a third evolution channel and account for some, at least, of the lowest luminosity R CrB and extreme helium stars. ",
        "conclusions": "Mergers of helium and carbon-oxygen white dwarfs are currently regarded to be the most favoured model for the origin of most R\\,CrB and EHe stars. The merger of two helium white dwarfs has been widely considered responsible for the origin of some hot subdwarfs, particulalry those with hydrogen-deficient surfaces. In this paper we have demonstrated that the most massive helium+helium white dwarf mergers can produce hot subdwarfs which subsequently become cool supergiants and, consequently, provide an alternative means to produce low-mass R\\,CrB and EHe stars. However, the number of R\\,CrB stars produced in this way may be  some 14 to 70 times smaller than from the CO+He WD merger channel. Nucleosynthesis of elements during the hot (fast-accretion) phase of the merger is roughly consistent with the observed abundances of  \\iso{12}{C}, \\iso{18}{O}, \\iso{19}{F} and \\iso{22}{Ne} in  R\\,CrB and EHe stars. Further work is required to establish whether the products of massive He+He WD mergers can be clearly distinguished from the products of CO+WD mergers."
    },
    "1208/1208.6149_arXiv.txt": {
        "abstract": "{Most observational studies so far point towards brown dwarfs sharing a similar formation mechanism as the one accepted for low mass stars. However, larger databases and more systematic studies are needed before strong conclusions can be reached. } {In this second paper of a series devoted to the study of the spectroscopic properties of the members of the Lambda Orionis Star Forming Region, we study accretion, activity and rotation for a wide set of spectroscopically confirmed members of the central star cluster Collinder 69 to draw analogies and/or differences between the brown dwarf and stellar populations of this cluster. Moreover, we present comparisons with other star forming regions of similar and different ages to address environmental effects on our conclusions.} {We study prominent photospheric lines to derive rotational velocities and emission lines to distinguish between accretion processes and chromospheric activity. In addition, we include information about disk presence and X-ray emission. } {We report very large differences in the disk fractions of low mass stars and brown dwarfs ($\\sim$58\\%) when compared to higher mass stars  (26$^{+4}_{-3}$ \\%) with 0.6 M$_{\\odot}$ being the critical mass we find for this dichotomy. As a byproduct, we address the implications of the spatial distribution of disk and diskless members in the formation scenario of the cluster itself. We have used the H$\\alpha$ emission to discriminate among accreting and non-accreting sources finding that  38$^{+8}_{-7}$ \\% of sources harboring disks undergo active accretion and that his percentage stays similar in the substellar regime. For those sources we have estimated accretion rates. Finally, regarding rotational velocities, we find a high dispersion in v$\\sin(i)$ which is even larger among the diskless population. } {} ",
        "introduction": "This is the second paper of a series devoted to studying, from a spectroscopic point of view, the young population present in several associations belonging to the Lambda Orionis Star Forming Region (LOSFR). In this paper, we systematically analyze several properties of the confirmed members of the central cluster Collinder 69 (C69); including the presence of disks, accretion onto the central star/brown dwarf, rotation and activity.  In \\cite{Bayo11} (from now on Paper I), we presented a very detailed and complete spectroscopic census of Collinder 69; the oldest ($\\sim$ 5 -- 12 Myr) of the associations belonging to the LOSFR. In short, this star forming region is located at $\\sim$400 pc \\citep{Murdin77} representing the head of the Orion giant. Its center is dominated by the O8III multiple star $\\lambda$ Orionis and comprises both recently formed stars from 0.2 M$_{\\odot}$ to 24 M$_{\\odot}$ and dark clouds actively forming stars.  The main goals of this paper are to study in detail the properties of the C69 confirmed members (e.g. rotation, activity, disk fractions and accretion rates), analyze the different populations within the cluster, and compare our results with other low mass star forming regions of similar and different ages. In fact, the evolutionary status of C69 seems to be specially suited to the study of disk evolution and accretion at the low end of the mass spectrum. According to our observations, in almost every bin in mass (for masses lower or equal to 0.6M$_{\\odot}$, $\\sim$3700K assuming a 5 Myr isochrone from \\citealt{Baraffe98}), we find a diversity of disk harboring sources: from those with optically thick disks undergoing active accretion (onto the central star) to others that have completely lost their primordial circumstellar disks. Finally, and as a byproduct of this study, we try to put our findings in context of the current theory of how the LOSFR as a whole was formed (triggered by a supernova, see \\citealt{DM01}, and more discussion on this scenario on Paper I).  This paper is organized as follows: In Section~\\ref{sec:data} we provide a description of the data analyzed. In Sections~\\ref{sec:rotvel} to ~\\ref{sec:distrib} properties such as rotation velocity, activity levels, accretion processes, disks, variability and spatial distribution of the population of C69 are studied. And, finally, our summary conclusions are presented in Section~\\ref{sec:conclusions}.  Unlike in Paper I, we have grouped the most interesting/puzzling objects into different categories (following the subsections of Sections~\\ref{sec:rotvel} and~\\ref{sec:acacc}) and discuss their peculiarities in Appendix A. ",
        "conclusions": ""
    },
    "1208/1208.6180_arXiv.txt": {
        "abstract": "{Aperture array (AA) technology is at the forefront of new developments and discoveries in radio astronomy. Currently LOFAR is successfully demonstrating the capabilities of dense and sparse AA's at low frequencies. For the mid-frequencies, from 450 to 1450\\,MHz, AA's still have to prove their scientific value with respect to the existing dish technology. Their large field-of-view and high flexibility puts them in an excellent position to do so. The Aperture Array Verification Program is dedicated to demonstrate the feasibility of AA's for science in general and SKA in particular. For the mid-frequency range this has lead to the development of EMBRACE, which has already demonstrated the enormous flexibility of AA systems by observing \\HI\\ and a pulsar simultaneously. It also serves as a testbed to demonstrate the technological reliability and stability of AA's. The next step will put AA technology at a level where it can be used for cutting-edge science. In this paper we discuss the developments to move AA technology from an engineering activity to a fully science capable instrument. We present current results from EMBRACE, ongoing tests of the system, and plans for EMMA, the next step in mid-frequency AA technology.}  \\FullConference{Resolving the Sky - Radio Interferometry: Past, Present and Future -RTS2012\\\\ April 17-20, 2012\\\\ Manchester, UK}   \\begin{document}  \\begin{figure} \\centering \\includegraphics[width=12cm]{embrace.png} \\caption{Actual image of a part of one of the EMBRACE stations.} \\label{embrace} \\end{figure} ",
        "introduction": "Aperture arrays (AA's, see Fig.~\\ref{embrace}) have been under development at ASTRON for decades, and the technology is now reaching a readiness level at which it can be applied in the construction of future scientific instruments, most notably the Square Kilometre Array (SKA). In this paper we discuss the so-called 'dense' aperture arrays, where the individual receiver elements are closely spaced, and the system is optimized to observe at a frequency range from roughly 450 to 1450\\,MHz. This technology is part of the Advanced Instrumentation Programme for the SKA.  Dense AA's have many advantages over dishes. The individual elements are intrinsically broad band receivers, e.g. for Vivaldi elements a frequency range of 9:1 is easily achievable, therefore these are a strong candidate for the individual antenna elements, see also \\cite{andy_talk}. Due to computational limitations at these higher frequencies, the operational range required is set to a more feasible 3:1. Since AA's have no moving parts, they can be repointed in a matter of seconds. This makes the technology extremely suitable for rapid follow-up of transient events.  A unique feature of AA's is the possibility to observe in multiple directions simultaneously (see also Fig.~\\ref{emma_multibeam}). With the field-of-view in each direction being an order of magnitude larger than that for dishes with single pixel feeds, AA's form an ideal technology to build a cheap and highly efficient survey instrument. The size of the field-of-view is in principle only limited by the available processing power, creating the potential of unprecedented survey speeds. In order to exploit this, cost reduction of beam forming and post-processing is the main challenge.  AA's employ multi-level beam forming. First, the wide beams of individual receiver elements are combined into a smaller beam, which is the equivalent of a primary beam in a dish. This defines the field-of-view of the instrument, and we will refer to it as such. Subsequently, all available fields are digitally combined to achieve the highest possible spatial resolution. We will refer to this as a beam. ",
        "conclusions": "The testing and verification of EMBRACE has delivered valuable insights, as well as ground-breaking results. We have successfully observed simultaneously in two independent directions, and achieved the required high accuracy in beam forming and calibration. With both stations operational, additional tests will continue throughout this year and into the next. In parallel, the development of improved tiles will take place for the next step: EMMA.  EMMA will be a science capable AA instrument, focussing on detecting the BAO signal and demonstrating imaging capabilities. Though it will be small, it can match survey speed with most of the upcoming and existing radio observatories in the mid-frequencies. Its multiple, large fields-of-view help to overcome the limited point-source sensitivity. Most notably, EMMA will provide access to a unique and continuous frequency range, which is only marginally covered by existing observatories. With the current specifications, EMMA will be a powerful and very flexible system, capable of a large range of scientific observations. It will make important progress in our understanding of the use of AA's in a scientific instrument, a necessary step in the evolution of AA technology towards the SKA.  \\subsection*{Acknowledgements} The authors would like to thank all the people involved in the construction, testing and operation of EMBRACE. Most notably, we thank Steve Torchinsky for sharing with us the results of the Nan\\c{c}ay EMBRACE station. Thanks is also due to the people in the astronomy group at ASTRON for extensive discussion on the capabilities of EMMA. We thank Russ Scott at Swinburne Astronomy Productions for a very pleasant collaboration in making the animation and images."
    },
    "1208/1208.1844_arXiv.txt": {
        "abstract": "High speed photometric observations of the spectroscopically-discovered PG 1159 star SDSS\\,J034917.41-005917.9 in 2007 and 2009 reveal a suite of pulsation frequencies in the range of 1038 -- 3323 $\\mu$Hz with amplitudes between 3.5 and 18.6 mmag. SDSS\\,J034917.41-005917.9 is therefore a member of the GW Vir class of pulsating pre-white dwarfs. We have identified 10 independent pulsation frequencies that can be fitted by an asymptotic model with a contant period spacing of 23.61 $\\pm$ 0.21 s, presumably associated with a sequence of $\\ell = 1$ modes. The highest amplitude peak in the suite of frequencies shows evidence for a triplet structure, with a frequency separation of 14.4 $\\mu$Hz. Five of the identified frequencies do not fit the $\\ell = 1$ sequence, but are, however, well-modeled by an independent asymptotic sequence with a constant period spacing of 11.66 $\\pm$ 0.13 s. It is unclear to which $\\ell$ mode these frequencies belong. ",
        "introduction": "The GW Vir stars are a sub-group of variables in the spectroscopic PG 1159 class, which form a link between the (post-AGB) central stars of planetary nebulae and the H-deficient white dwarf cooling sequence. They pulsate non-radially and lie in an instability strip bounded by effective temperatures $200\\,000 > T_{\\rm eff} > 75\\,000$ K, excited by the kappa mechanism working through partial ionization of carbon and oxygen. Studying these stars with astroseismology has provided important knowledge on the interiors of the late stages of stellar evolution (Winget \\& Kepler 2008). There are 19 known GW Vir stars (Quirion, Fontaine \\& Brassard 2007; Quirion 2009), showing a wide variety of behaviour. The possible addition of more examples is therefore of significance. Here we show that the known spectroscopic PG 1159 star SDSS J034917.41-005919.2 (hereafter SDSS J0349-0059) is a non-radial pulsator, putting it in the GW Vir subclass. In Sect.~2 we describe what is already known about SDSS J0349-0059 and list our high speed photometric observations. Sect.~3 analyses these and presents comparisons with other GW Vir stars. ",
        "conclusions": "Our photometric observations show that SDSS\\,J0349-0059 is a non-radial pulsator in the GW Vir class of variable stars. Frequency splitting of the principal oscillation mode at 419 s reveals a rotation period of 0.40 d, if the oscillation mode is $\\ell = 1$. As with many other GW Vir stars, it is possible to represent most of the modes with a linear relationship in period, similar to what has been seen in PG\\,1707+427, but with parameters that put SDSS J0349 slightly redward of the latter star, which now defines the red edge of the GW Vir instability strip.  The five oscillation modes not included in the above group can be fitted by another linear relationship in which the spacing is a harmonic of the first sequence. However, there is no currently known physical model which explains this behaviour. This is unlike PG\\,1707+427 where the discrepant oscillation modes appear to have an $\\ell = 2$ origin."
    },
    "1208/1208.6419_arXiv.txt": {
        "abstract": "In this paper, we present photometry for young star clusters in M31, which are selected from Caldwell et al. These star clusters have been observed as part of the Beijing--Arizona--Taiwan--Connecticut (BATC) Multicolor Sky Survey from 1995 February to 2008 March. The BATC images including these star clusters are taken with 15 intermediate-band filters covering 3000--10000 \\AA. Combined with photometry in the {\\sl GALEX} far- and near-ultraviolet, broad-band $UBVRI$, SDSS $ugriz$, and infrared $JHK_{\\rm s}$ of Two Micron All Sky Survey, we obtain their accurate spectral energy distributions (SEDs) from $1538-20000$ \\AA. We derive these star clusters' ages and masses by comparing their SEDs with stellar population synthesis models. Our results are in good agreement with previous determinations. The mean value of age and mass of young clusters ($<2$ Gyr) is about 385 Myr and $2\\times 10^4~{M_\\odot}$, respectively. There are two distinct peaks in the age distribution, a highest peak at age $\\sim$ 60 Myr and a secondary peak around 250 Myr, while the mass distribution shows a single peak around $10^4~{M_\\odot}$. A few young star clusters have two-body relaxation times greater than their ages, indicating that those clusters have not been well dynamically relaxed and therefore have not established the thermal equilibrium. There are several regions showing aggregations of young star clusters around the 10 kpc ring and the outer ring, indicating that the distribution of the young star clusters is well correlated with M31's star-forming regions. The young massive star clusters (age $\\leq 100$ Myr and mass $\\geq 10^4~{M_\\odot}$) show apparent concentration around the ring splitting region, suggesting a recent passage of a satellite galaxy (M32) through M31 disk. ",
        "introduction": "Star clusters are considered as important tracers for understanding the formation and evolution of their host galaxies \\citep{san10}. Star cluster systems have been traditionally separated into two populations--globular clusters and open clusters (GCs and OCs)--on their ages, masses, metallicities, and positions. However, more recent studies have discovered that the distinction between GCs and OCs becomes increasingly blurred \\citep[see][for details]{perina10}.  \\citet{gk52} listed photometric colors and magnitudes for star clusters in Magellanic Clouds (MCs) and the Fornax dwarf system and divided them into two groups. They found that star clusters in blue group have central condensation properties similar to those of the red group, which were considered as GCs, however they could not be identified with the Galactic OCs. \\citet{hodge61} termed 23 clusters in the Large Magellanic Cloud (LMC)--differing from GCs in their relative youth and OCs in their richness and shape--as ``young populous clusters'', which were called ``young massive clusters'' (YMCs) or ``blue luminous compact clusters'' (BLCCs) by \\citet{Fusi05}. Actually, the blue integrated colors for a cluster may be influenced by several factors , such as poor metallicity (the luminosity of the horizontal branch), young age (the position of the main-sequence turnoff stars), and some exotic stellar populations (e.g., blue stragglers, Wolf-Rayet stars). However, several studies \\citep[e.g.,][]{Williams01,Beasley04} have reached similar conclusions that the exceedingly blue colors of BLCCs are a direct consequence of their young ages \\citep[see][for details]{Fusi05}.  M31 is the largest galaxy in the Local Group, and has a large number of star clusters, including young clusters having been studied by many authors. \\citet{bohlin88,bohlin93} listed 11 objects in M31 classified as blue clusters using the UV colors, most of which have been proved to be young clusters \\citep{Fusi05,cald09,perina09,perina10}, except for B133 and B145, which were stated as a star and an old GC \\citep{cald09}, respectively. \\citet{cald09, cald11} derived ages and masses for a large sample of young clusters, and found that these star clusters are less than 2 Gyr old, and most of them have ages between $10^8$ and $10^9$ yr and masses ranging from $2.5 \\times 10^2~{M_\\odot}$ to $1.5 \\times 10^5~{M_\\odot}$. These authors also stated that the young star clusters in M31 show a range of structures, most of which have low concentrations typical of OCs in the Milky Way (MW), however, there are a few with high concentrations similar to the MW GCs. \\citet{vanse09} carried out a survey of compact star clusters in the southwest part of M31, and suggested a rich intermediate-mass star cluster population in M31, with a typical age range of 30 Myr $-$ 3 Gyr, peaking at $\\sim$ 70 Myr. In order to ascertain the properties of the BLCCs, \\citet{perina09,perina10} performed an image survey for 20 BLCCs lying in the disk of M31 using the Wide Field and Planetary Camera-2 (WFPC2) on the {\\it Hubble Space Telescope} ({\\it HST}). In addition, another key aim of this {\\it HST} survey was to determine the fraction of contamination of BLCCs by asterisms, since \\citet{Cohen05} suggested that a large fraction of the putative BLCCs may in fact be just asterisms. \\citet{Cohen05} presented the resulting $K'$ images of six very young or young star clusters in M31 observed with the Keck laser guide star adaptive optics system, and indicated that the four youngest out of these six objects are asterisms. However, \\citet{cald09} presented a conclusion that these four objects are true clusters based on spectra. The {\\it HST} images \\citep{perina09,perina10} showed that nineteen of the twenty surveyed candidates are real star clusters, and one (NB67) is a bright star. \\citet{barmby09} measured surface brightness profiles for 23 bright young star clusters using images from the WFPC2, including the sample clusters of \\citet{perina09,perina10}, and derived the structural properties by fitting the surface brightness profiles to several structural models. The authors stated that the sample young clusters are expected to dissolve within a few Gyr and will not survive to become old GCs, and that young star clusters in M31 and MCs follow the same fundamental plane relations as old GCs of M31, MCs, the MW and NGC 5128, regardless of their host galaxy environments. \\citet{johnson12} presented a M31 stellar cluster catalog utilizing the Panchromatic Hubble Andromeda Treasury survey data, which will cover $\\sim1/3$ of M31 disk with multiple filters and allow the identification of thousands of star clusters.  The large population of young star clusters reflect a high efficiency of cluster formation, possibly triggered by a current interaction event between M31 and its satellite galaxy \\citep{Gordon06, block06}, suggesting that young star clusters should be associated with the star-forming (SF) regions of M31. \\citet{fan10} found that young clusters ($<$ 2 Gyr) are spatially coincident with M31's disk, including the 10 kpc ring and the outer ring \\citep{Gordon06}. Although these authors also found the young star clusters in the halo of M31, all of the clusters outside of the optical disk of M31 are old, globular clusters \\citep[see][for details]{perina11}. \\citet{kang12} stated that most of young star clusters' kinematics have the thin, rotating disk component \\citep[see also][]{rey07}. The young star clusters' distribution has a distinct peak around $10-12$ kpc from the center in M31 disk, and some young star clusters show concentration around the 10 kpc ring splitting regions near M32 and most of them have systematically younger ages ($< 100$ Myr). \\citet{kang12} also stated that the young star clusters show a spatial distribution similar to OB stars, UV SF regions, and dust, all of which are important tracers of disk structures.  Several criteria were developed for selecting young clusters from the integrated spectrum and colors. \\citet{Fusi05} comprehensively studied the properties of 67 very blue and likely YMCs in M31 selected according to their color $[(B-V)_0\\leq 0.45]$ and/or the strength of $H\\beta$ spectral index ($\\rm H\\beta \\geq 3.5$ \\AA). \\citet{Peacock10} presented a catalog of M31 GCs based on images from the SDSS and the Wide Field CAMera on the United Kingdom Infrared Telescope and selected a population of young clusters with a definition of $[(g-r)_0<0.3]$. \\citet{kang12} published a catalog of M31 young clusters ($\\leq 1$ Gyr) and supported the selection criteria $[(NUV-r)_0\\leq 2.5]$ and $[(FUV-r)_0\\leq 3.0]$ \\citep{bohlin93,rey07}. These criterions may play important roles in distinguishing young from old clusters for those whose ages cannot be derived accurately.  The formation and disruption of young star clusters represent a latter-day example of the hierarchical formation of galaxies \\citep{fall04}. Motivated by that, we decided to describe some basic properties of young star clusters in M31, such as positions, distributions of ages and masses, correlations of the ages and masses with structure parameters, which may provide important information about the processes involved in their formation and disruption.  In this paper, we will provide photometry of a set of young star clusters in M31 using images obtained with the Beijing--Arizona--Taiwan--Connecticut (BATC) Multicolor Sky Survey Telescope. By comparing the observed  SEDs with the {\\sc galev} simple stellar population (SSP) models, we derive their ages and masses. This paper is organized as follows. In Section 2 we present the BATC observations of the sample clusters, the relevant data-processing steps, and the {\\sl GALEX} (FUV and NUV), optical broad-band, SDSS $ugriz$ and 2MASS NIR data that are subsequently used in our analysis. In Section 3 we derive ages and masses of the sample clusters. A discussion on the sample young clusters ($<$ 2 Gyr) will be given in Section 4. Finally, we will summarize our results in Section 5. ",
        "conclusions": "\\subsection{Position}  Figure 6 shows the number, ages and masses of young star clusters ($<2$ Gyr) as a function of projected radius from the center of M31, adopted at $\\rm \\alpha_0=00^h42^m44^s.30$ and $\\rm \\delta_0=+41^o16'09''.0$ (J2000.0) following \\citet{hbk91} and \\citet{per02}. In the top panel, the histogram for the radial distribution of young star clusters shows clearly two peaks at $4-7$ kpc and $9-11$ kpc, while in the middle panel and bottom panel, wide age and mass distributions can be seen in these two peak regions.  \\begin{figure} \\figurenum{6} \\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{fig6.eps}}} \\caption{ ($Top~panel$) Number histogram of young star clusters against projected radius. ($Middle~panel$) Age versus projected radius for sample young clusters. ($Bottom~panel$) Mass versus projected radius for sample young star clusters. The open rectangles show the two peaks around $4-7$ kpc and $9-11$ kpc of the radial distribution for young star clusters.} \\label{fig:fig6} \\end{figure}  \\citet{kang12} presented the radial distribution of clusters against the distance from the center of M31, and found that the young clusters show two peaks around $10-12$ kpc and $13-14$ kpc. They also found that the UV SF regions show two distinct peaks: a main peak at $\\sim16$ kpc and a secondary peak around 11 kpc. In addition, a small peak at $5-8$ kpc in the distribution of ages of UV SF regions against the projected radius \\citep[see Figure 19 of][]{kang12} can be clearly found. We argued that the peak at $4-7$ kpc obtained in this paper should be associated with the peak at $5-8$ kpc for the UV SF regions, while the peak at $9-11$ kpc obtained in this paper correlate with the well-known 10 kpc ring \\citep{Gordon06}.  Figure 7 displays the spatial distribution and radial distribution of the M31 young clusters with different age bins: (a) $t<0.1$ Gyr; (b) 0.1 Gyr $\\leq t<$ 0.4 Gyr; (c) 0.4 Gyr $\\leq t<$ 1 Gyr; (d) 1 Gyr $\\leq t<$ 2 Gyr. In the top panel, young star clusters in different age ranges are drawn with different marks. The inner, solid ellipse and the dashed contour represent the 10 kpc ring and the outer ring from \\citet{Gordon06} based on infrared observations with the Multiband Imaging Photometer for Spitzer (MIPS) instrument on the {\\it Spitzer Space Telescope}, respectively. The 10 kpc ring was drawn with a center offset from the M31 nucleus by [$5'.5$, $3'.0$] \\citep{Gordon06} with a radius of 44 arcmin (10 kpc). There are several regions drawn with open rectangles which show aggregations of young star clusters to different extents. However, we should point out that these aggregations of young star clusters may be caused by the projection effect because of the inclination of M31 disk. \\citet{vanse09} noted two clumps of young clusters, both of which are located in one rectangle ($\\sim-13$ kpc $< X <$ $-9$ kpc and $-3$ kpc $< Y <$ 0 kpc). The star clusters in this study are spatially coincident with the disk and the rings, indicating that the distribution of the young star clusters correlates with the galaxy's SF regions, which is consistent with previous studies \\citep{fan10,kang12}. In the bottom panel, the number of young star clusters in different age bins as a function of projected radial distance from the M31 center was shown. We can see that clusters younger than $0.1$ Gyr show most obvious aggregation around the 10 kpc ring.  \\begin{figure} \\figurenum{7} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{fig7.eps}}} \\caption{Spatial distribution ($top~panel$) and radial distribution ($bottom~panel$) of M31 young star clusters with different age bins: (a) $t<$ 0.1 Gyr; (b) 0.1 Gyr $\\leq t<$ 0.4 Gyr; (c) 0.4 Gyr $\\leq t<$ 1 Gyr; (d) 1 Gyr $\\leq t<$ 2 Gyr. The inner, solid ellipse and the dashed contour represent the 10 kpc ring and the outer ring from \\citet{Gordon06}, while the dotted ellipse is the M31 disk/halo boundary as defined by \\citet{rac91}. The several small rectangles show the clumps of young clusters to the extents.} \\label{fig:fig7} \\end{figure}  \\citet{Gordon06} ran a number of numerical simulations of the M31--M32 and M31-NGC 205 interactions, and assumed a passage of M32 through the disk of M31 occurring 20 Myr ago, resulting in a burst of star formation that propagates outward through the disk. \\citet{block06} suggested that M32 and M31 had an almost head-on collision about 210 Myr ago, and M32 passed through M31 disk again about 110 Myr ago \\citep[see Figure 2 of][]{block06}, which induced two off-center rings--an inner ring with projected dimensions of $\\sim$ 1.5 kpc and the 10 kpc ring. Both of the simulations recurred the 10 kpc ring and the observed split.  We divided our sample star clusters younger than 300 Myr into six groups, and showed the spatial distribution for each group in Figure 8. We can see that only star clusters with ages 50 Myr $-$ 100 Myr appear around the 10 kpc ring and the ring splitting region ($-9.5$ kpc $< X <$ $-7.5$ kpc and $-2.5$ kpc $< Y <$ $-0.5$ kpc) \\citep{kang12}, indicating that 1) the 10 kpc ring may begin to form about 100 Myr ago; 2) M32 passed through the southern part of M31 disk around 100 Myr and in turn resulted in the split in the form of a hole. This appears to be consistent with the prediction by \\citet{block06} of a second passage of M32 about 110 Myr ago. After the second passage, star clusters formed around the split for a long period, since there are a number of star clusters around the split with ages younger than 50 Myr. \\citet{davidge12} reported that the star formation rate (SFR) of the M31 disk would be elevated greatly and quickly after an encounter event, and it would finally drop when the interstellar medium is depleted and disrupted. However, from Figure 8, we cannot find evidence of radial trend of star cluster ages \\citep[see also][]{kang12,cald09}.  \\begin{figure} \\figurenum{8} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{fig8.eps}}} \\caption{Spatial distribution of six groups of M31 young star clusters younger than 300 Myr, divided with same age bin of 50 Myr. The inner, solid ellipse and the dashed contour represent the 10 kpc ring and the outer ring from \\citet{Gordon06}, while the dotted ellipse is the M31 disk/halo boundary as defined by \\citet{rac91}. The small rectangle represents the ring splitting region in the southern part of M31 disk.} \\label{fig:fig8} \\end{figure}  Figure 9 shows the spatial and radial distribution of the M31 young star clusters with different mass bins: (a) $10^2~{M_\\odot} \\leq M < 10^3~{M_\\odot}$; (b) $10^3~{M_\\odot} \\leq M < 10^4~{M_\\odot}$; (c) $M \\geq 10^4~{M_\\odot}$. In the top panel, clusters of these three groups are drawn with different marks. The bottom panel presents the number of young clusters in different mass bins as a function of projected radial distance from M31 center, and it shows that young clusters more massive than $10^4~{M_\\odot}$ are most concentrated nearby the 10 kpc ring.  \\begin{figure} \\figurenum{9} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{fig9.eps}}} \\caption{Spatial distribution ($top~panel$) and radial distribution ($bottom~panel$) of M31 young star clusters with different mass bins: (a) $10^2~{M_\\odot}$ $\\leq M<$ $10^3~{M_\\odot}$; (b) $10^3~{M_\\odot}$ $\\leq M<$ $10^4~{M_\\odot}$; (c) $M>$ $10^4~{M_\\odot}$. The inner, solid ellipse and the dashed contour represent the 10 kpc ring and the outer ring from \\citet{Gordon06}, while the dotted ellipse is the M31 disk/halo boundary as defined by \\citet{rac91}. The several small rectangles show the clumps of young clusters to the extents.} \\label{fig:fig9} \\end{figure}  \\subsection{Age and Mass Distribution}  Figure 10 plots the distribution of estimated ages and masses for the young star clusters. A prominent correlation can be seen that mass increases with age. There are two distinct peaks in the age histogram: a highest peak at age $\\sim$ 60 Myr ($\\log \\rm age=7.8$) and a secondary peak around 250 Myr ($\\log \\rm age=8.4$). The mass distribution of the young star clusters show a single peak around $10^4~{M_\\odot}$. The mean values of age and mass of young clusters are about 385 Myr and $2\\times 10^4~{M_\\odot}$, slightly higher than the values presented by \\citet{kang12}, which are 300 Myr and $10^4~{M_\\odot}$, respectively. Most of our young clusters have masses ranging from $10^{3.5}~{M_\\odot}$ to $10^{5}~{M_\\odot}$, which are more massive than OCs in the solar neighborhood \\citep{piskunov08}, but less massive than typical GCs in the MW \\citep{mm05}. The lack of young clusters more massive than $10^5~{M_\\odot}$ is also noted by \\citet{vanse09} and \\citet{cald09}, possibly caused by a low-average SFR of M31 \\citep{barmby06} or hidden by dust clouds in the disk due to the inclination angle of M31 \\citep{vanse09}.  \\begin{figure} \\figurenum{10} \\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{fig10.eps}}} \\caption{Age and mass distribution of the sample young star clusters in this paper. The histograms for age and mass are presented with gray colors.} \\label{fig:fig10} \\end{figure}  \\citet{portegies10} have listed three phases for the evolution of a young star cluster: 1) the first few Myr, during which the star formation activity is still proceeding and the star cluster is rich in gas; 2) a subsequent period after the first supernovae (some 3 Myr after formation), in which a young cluster is experiencing a serious loss of gas and dust, and stellar mass loss plays an important role in the cluster evolution; 3) a later stage that stellar dynamical processes dominate the cluster evolution. The dividing line between phase 2 and phase 3 may be anywhere between 100 Myr and 1 Gyr, and most of our young clusters are experiencing the phase 2 or phase 3.  \\citet{cs87} presented that after 5 Gyr, both mass and galactic location are important evolutionary parameters for GCs. \\citet{spitzer58} discussed the destructive effects of encounters of clusters with giant molecular clouds (GMCs), and presented that the disruption time for a star cluster varies directly with the cluster density and is about 200 Myr for a mean density of 1 $M_\\odot/{\\rm pc}^3$. \\citet{sh58} also reported that two-body relaxation is effective at destroying low-mass clusters and this may account for the scarcity of low-mass older clusters. Actually, the two-body relaxation and the encounters with GMCs are also important processes that lead to young cluster disruption \\citep[see][and references therein]{cald09}, while \\citet{portegies10} presented that mass loss due to stellar evolution is the most important process in the young cluster dissolution. It is evident that star cluster mass is one key parameter in the star cluster evolution. \\citet{bl03} derived an empirical relation between the disruption time and the initial mass of star clusters in the solar neighborhood, Small Magellanic Cloud (SMC), M51, and M33. \\citet{lamers05} determined a disruption time of 1.3 Gyr for a $10^4~{M_\\odot}$ cluster in the solar neighborhood, while \\citet{cald09} reported that most of M31 young clusters would be destroyed in the next Gyr or so, and only some massive and dense ones may survive for a longer time.  Several features are shown in Figure 10: 1) there is an obvious gap in the age distribution around 100 Myr. 2) there are few clusters older than 400 Myr (${\\log \\rm age=8.6}$) with mass lower than $10^4~{M_\\odot}$. Although many low-mass clusters can be easy to disrupt, this gap may be caused by a selection effect. In fact, \\citet{johnson12} found that the completeness of M31 ground-based sample drops precipitously at $m_{\\rm F475W}>18$ ($M_{\\rm F475W}>-6.5$), which is about $2\\times10^4~{M_\\odot}$. 3) there are few clusters more massive than $10^5~{M_\\odot}$, which may be caused by a low-average SFR of M31 or the hidden by dust clouds in the M31 disk as discussed above. 4) there is a gap of clusters with very low masses ($\\sim10^3~{M_\\odot}$) and younger than 30 Myr (${\\log \\rm age=7.5}$). These clusters may be too faint to be sample objects of \\citet{cald09,cald11}, indicating that our sample is not complete in these age and mass ranges \\citep[see also][]{cald09}.  Figure 11 shows the age distribution in different mass intervals (top panel) and mass function in different age intervals (bottom panel). The histograms are derived using a 0.4-dex bin width with different starting values. These distributions contain information about the formation and disruption history of star clusters \\citep{fc12}, however, the interpretation of the empirical distributions of clusters depends strongly on how incompleteness affects the sample \\citep{gieles07}. In the top panel, we can see an obvious gap before 40 Myr ($\\log \\rm age=7.6$), which is caused by a selection effect. The age distribution of the clusters does not declines monotonically, with an apparent bend around 200 Myr ($\\log \\rm age=8.3$). We argued that this bend near 200 Myr may be explained as a burst of cluster formation, possibly caused by a current interaction event between M31 and its satellite galaxy, such as the collision between M31 and M32 about 210 Myr ago suggested by \\citet{block06}. The two decline trends starting from 40 Myr and 200 Myr reflect a rapid disruption of clusters. \\citet{vanse09} noted a peak of the cluster age distribution at 70 Myr, and suggested an enhanced cluster formation episode at that epoch. In the bottom panel, the gap in the number of clusters in the low-mass regions ($\\log \\rm mass < 3.5$) is apparently due to a sample incompleteness, but not physical \\citep{vanse09}. The initial mass function for star clusters should be slightly steeper \\citep{fc12} than what is shown here because of the short lifetimes of low-mass clusters. Recently, \\citet{fc12} compared the observed age distributions and mass functions of star clusters in the MW, MCs, M83, M51, and Antennae, and found that these distributions of clusters are similar in different galaxies. However, due to the incompleteness of our cluster sample, partly due to the exclusion of clusters that cannot derive accurate photometry, we would not give any empirical formulas of the distributions for age and mass.  \\begin{figure} \\figurenum{11} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{fig11.eps}}} \\caption{Age distribution of the sample young star clusters with differen mass intervals ($top~panel$) and mass function with differen age intervals ($bottom~panel$). The histograms are derived using a 0.4-dex bin width with different starting values.} \\label{fig:fig11} \\end{figure}  \\subsection{Correlations with Structure Parameters}  In this section, we will discuss the correlations of ages and masses with structure parameters, which are derived by King-model \\citep{king66} fits for clusters in M31 \\citep{bhh02, barmby07, barmby09}. Because the sample clusters are younger than 2 Gyr, the structure parameters obtained from the bluer filters are preferred \\citep[see][in detail]{barmby09}. There are four clusters (B315, B319, B368, and B374) which have been studied twice by \\citet{bhh02, barmby07} and \\citet{barmby09}, and we would use the new results in \\citet{barmby09}.  Figure 12 shows structure parameters as a function of age for young clusters in this paper. Some correlations can be seen, the concentration $c$, defined as $c\\equiv\\log(r_t/r_0)$, decreases with age. The trend is largely driven by clusters B342 and B368, both of which have large $c$ values (3.98 for B342 and 3.87 for B368). Both the scale radius $r_0$ and projected core radius $R_c$ increase with age. Clusters B342 and B368 have very small $r_0$ and $R_c$ values and are drawn with arrows in Figure 12 (The values of $r_0$ and $R_c$ are $\\sim$ 0.014 pc for B342, while are $\\sim$ 0.011 pc for B368). \\citet{elson89} and \\citet{elson91} discussed the trend for core radius against age, and argued that this trend may represent real evolution in the structure of clusters as they grow old, partially explained by the effect of mass segregation \\citep{mg03}, or dynamical effects such as heating by black hole (BH) binaries \\citep{mackey07}. \\citet{wilkinson03} also demonstrated that neither large differences in primordial binary fraction nor a tidal heating due to differences in the cluster orbits could account for the observed trend. The best-fit central surface brightness $\\mu_{V,0}$ shows a decreasing trend with age, and \\citet{barmby09} argued that this trend may be likely due to the fading of stellar population and the increase of core radius $R_c$ with age. We also see that the central mass density $\\rho_0$ decreases with age, although the scatters are great. \\citet{barmby09} presented that the central mass density shows very little trend with age for both the M31 young clusters and young clusters in the MCs. There is no obvious correlation between $t_{r,h}$, the two-body relaxation time at the model-projected half-mass radius, and age. The dashed line represents the region that $t_{r,h}$ equal to age. It can be seen that most clusters (except for DAO38 and M091) have ages less than $t_{r,h}$, indicating that these young clusters have not been well dynamically relaxed. Because two-body encounters can transfer energy between individual stars and then impel the system to establish thermal equilibrium \\citep{portegies10}, we argue that these young clusters have not established thermal equilibrium.  \\begin{figure} \\figurenum{12} \\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{fig12.eps}}} \\caption{Structure parameters as a function of age for the sample young star clusters in this paper.} \\label{fig:fig12} \\end{figure}  Figure 13 shows structure parameters as a function of mass for the sample young clusters. The concentration $c$ increases with mass, although the trend is much weak. \\citet{fc12} presented $c$ plotted against mass for clusters in MCs, and found that there was no correlation between $c$ and mass. Actually, we found that all the clusters in \\citet{fc12} have $c$ less than 2.5, much smaller than the largest value in our sample ($\\sim4$). If we do not include the two clusters B342 and B368, the correlation for $c$ with mass nearly disappear. Both $r_0$ and $R_c$ increase with mass, however, the trend is largely weaken by cluster B327, which has very small values of $r_0$ and $R_c$, but larger than those of B342 and B368 which are drawn with arrows in Figure 13. Both the central surface brightness $\\mu_{V,0}$ and central mass density $\\rho_0$ decrease weakly with mass, while no obvious correlation between $t_{r,h}$ and mass can be seen.  \\begin{figure} \\figurenum{13} \\resizebox{\\hsize}{!}{\\rotatebox{-90}{\\includegraphics{fig13.eps}}} \\caption{Structure parameters as a function of mass for the sample young star clusters in this paper.} \\label{fig:fig13} \\end{figure}  We checked the surface brightness profiles of B327, B342, and B368 displayed in \\citet{barmby09}, which have very small $r_0$ and $R_c$ and very large $\\mu_{V,0}$ and  $\\rho_0$, and found that these core profiles are cuspy. \\citet{barmby09} concluded that the cores of these clusters did not appear to be resolved in the {\\it HST}/WFPC2 images and the structural parameters for these clusters would be uncertain if the central cluster luminosity is dominated by only a few bright stars. However, if these cuspy core profiles are true integrated properties, which may be better fitted by a power-law structure model \\citep[e.g.,][]{sersic68}, the three clusters may have been post core-collapse \\citep[see][in detail]{tanvir12}.  \\subsection{Young Massive Clusters}  YMCs are often related to the violent SF episodes triggered by galaxy collisions, mergers, and close encounters \\citep{grijs07}. However, based on a sample of 21 nearby spirals, \\citet{lr99} found that YMCs can exist in a wide variety of host galaxy environments, including quiescent galaxies, and that there is no correlation between the morphological type of the galaxies and their contents of YMCs. YMCs are dense aggregates of young stars, which are also expected to be the nurseries for many unusual objects, including exotic stars, binaries, and BHs \\citep{portegies10}. Many studies \\citep{barmby09, cald09, vanse09, Peacock10, perina10, portegies10, ma11} that focused on M31 YMCs have derived remarkable achievements in understanding their stellar populations, structure parameters, and dynamical properties.  There are 13 YMCs in our cluster sample with a definition of age $\\leq 100$ Myr and mass $\\geq 10^4~{M_\\odot}$ \\citep{portegies10}. Figure 14 shows the spatial distribution of the 13 YMCs, while different sizes of the open circles indicate YMCs in different mass ranges. The rectangle between the 10 kpc ring and the outer ring represents the split in the southern part of the M31 disk, and the two black filled triangles represent M32 and NGC 205. It is not surprising to see that most of the YMCs gather around the split, indicating that there has been a high-level star formation activity, which is consistent with previous studies \\citep{Gordon06, kang12}.  \\begin{figure} \\figurenum{14} \\resizebox{\\hsize}{!}{\\rotatebox{0}{\\includegraphics{fig14.eps}}} \\caption{Spatial distribution for YMCs drawn with different sizes of the open circles indicating different mass ranges. The inner, solid ellipse and the dashed contour represent the 10 kpc ring and the outer ring from \\citet{Gordon06}, while the dotted ellipse is the M31 disk/halo boundary as defined by \\citet{rac91}. The small rectangle represents the ring splitting region in the southern part of M31 disk, and the two filled black triangles represent M32 and NGC 205.} \\label{fig:fig14} \\end{figure}"
    },
    "1208/1208.3650_arXiv.txt": {
        "abstract": "{About 10\\,\\% of white dwarfs have magnetic fields with strength in the range between about $10^5$ and $5\\, 10^8$~G. It is not known whether the remaining white dwarfs are not magnetic, or if they have magnetic fields too weak to be detected with the techniques adopted in the large surveys. Information is particularly lacking for the cooler (and generally fainter) white dwarfs.} {We describe the results of the first survey specifically devised to clarify the detection frequency of kG-level magnetic fields in cool DA white dwarfs.} {Using the FORS1 instrument of the ESO VLT, we have obtained Balmer line circular spectropolarimetric measurements of a small sample of cool (DA6 -- DA8) white dwarfs.  Using FORS and UVES archive data, we have also revised numerous white dwarf field measurements previously published in the literature.} { We have discovered an apparently constant longitudinal magnetic field of $\\sim 9.5$\\,kG in the DA6 white dwarf \\object{WD\\,2105$-$820}. This star is the first weak-field white dwarf that has been observed sufficiently to roughly determine the characteristics of its field.  The available data are consistent with a simple dipolar morphology with magnetic axis nearly parallel to the rotation axis, and a polar strength of $\\simeq 56$\\,kG.  Our re-evaluation of the FORS archive data for white dwarfs indicates that longitudinal magnetic fields weaker than 10\\,kG have previously been correctly identified in at least three white dwarfs. However, for one of these three weak-field stars (\\object{WD\\,2359$-$434}), UVES archive data show a $\\sim 100$\\,kG mean field modulus.  Either at the time of the FORS observations the star's magnetic field axis was nearly perpendicular to the line of sight, or the star's magnetic field has rather complex structure.  } {We find that the probability of detecting a field of kG strength in a DA white dwarf is of the order of 10\\,\\% for each of the cool and hot DA stars. If there is a lower cutoff to field strength in white dwarfs, or a field below which all white dwarfs are magnetic, the current precision of measurements is not yet sufficient to reveal it.} ",
        "introduction": "In 1970, a magnetic field was discovered in the peculiar white dwarf (WD) Grw\\,$+70\\,8247$ = GJ\\,472 \\citep{Kempetal70}. The field strength was eventually estimated to be of the order of 300\\,MG \\citep{Greenstein84,WickramasingheFerrario88,Jordan92}. Since this first detection of a magnetic field in a degenerate star, about 200 magnetic white dwarfs (MWDs) have been discovered \\citep{Kawkaetal07,Kulebietal09b}. It is found that about 10\\,\\% of all single WDs have a magnetic field with a strength in the range between hundreds of kG and hundreds of MG.  It is not at all clear how the magnetic fields in WDs originate, nor what information they carry about the origin and evolution of magnetism during stellar evolution. It is also not very clear yet how these fields influence such phenomena as rotation periods or pulsation of white dwarfs. Clearly, a broad observational base of data is essential for understanding these issues.  The magnetic fields of WDs are sometimes variable with the stellar rotation period, which when measurable is typically of the order of hours or days \\citep[e.g. ][]{Kawkaetal07}. It appears that MWDs may often be somewhat more massive than the overall WD average mass of about $0.6\\,M_\\odot$ \\citep{Liebert88}, although fields are occasionally found in relatively low-mass WDs. Most of the fields known are in WDs of spectral type DA, a white dwarf classification indicating that the optical spectrum shows only spectral lines of hydrogen, and which generally identifies WDs with H-rich atmospheres. This is at least partly a selection effect due to the fact that the strong and sharp Balmer lines are particularly sensitive probes of stellar magnetism, which in many cases can be easily detected in low-dispersion spectra from surveys such as the Sloan Digital Sky Survey \\citep{Kulebietal09a,Kulebietal09b}.  The concentration of WD magnetic field strengths as a function of $\\log \\bs$ \\citep{Kulebietal09b} in the best-studied range of 1--100\\,MG has raised the question of whether there is a cutoff field strength below which white dwarf fields do not occur \\citep[as is the case for Ap stars,][]{Auriereetal07}, or whether the probability of detecting a field might rise sharply below a field strength of some tens of kG.  Resolving this question confronts the difficulty of detecting weak fields in such faint, broad-lined objects, and our current knowledge of the low-field tail of the white dwarf field strength distribution is limited mainly by instrumental constraints. It is very difficult to obtain field measurements with standard errors of less that about 10~kG without using the largest available telescopes \\citep[see e.g.][]{Valyavinetal06, Kawkaetal07}. However, the study of available statistics by \\citet{Liebertetal03} and the survey by \\citet{Aznaretal04} both suggest that the detection rate for field weaker than a few tens of kG may be significantly higher than the frequency of $\\sim 10$\\,\\%, which characterises the overall detection rate of stronger fields.  A further question of great interest is whether the magnetic fields of WDs evolve with time, and if so, how they evolve. The searches for kG fields reported so far \\citep{Aznaretal04,Valyavinetal06,Kawkaetal07,Jordanetal07} have focussed almost entirely on the generally brighter hotter (and therefore younger) white dwarfs. It is thus worthwhile to focus a survey on cooler and older white dwarfs, and the higher detection probability predicted by earlier work suggests that even a fairly small sample of such stars may yield interesting results.  Thus, to increase the available information about the incidence of weak fields, and to extend this information to include some older, cooler white dwarfs, we have carried out a modest survey for fields in DA WDs with effective temperatures \\te\\ below about 14\\,000~K, aiming at obtaining field measurements with $\\sim 1$\\,kG error bars.  Recent work by \\citet{Bagnuloetal12}, \\citet{Jordanetal12}, and \\citet{Lanetal12} have shown that the results of some FORS1 surveys of magnetic fields in various classes of stars were affected by spurious detections, highlighting the need for a re-analysis of published data for MWDs. Therefore, we have complemented the results of our own survey with the revision of all FORS1 field measurements of WDs. ",
        "conclusions": "The database that we have considered includes 20 hot DA stars (generally spectral type DA1 to DA4, $\\te \\ga 14000$\\,K) and 15 cool DA stars (spectral type DA5 to DA8; $\\te \\la 14000$\\,K). (We omit 40 Eri B from our sample, as we have no data to confirm the field detected, and the star was not observed in a survey of known size.) Since there are two firmly detected MWDs in each of the hot and cool samples, we conclude that detection rates are about 10\\,\\% for the hot sample, and 13\\,\\% for the cool sample. The small size of the sample and the small number of detections set a serious limit to accuracy of these frequency estimates. Using the Wilson 95\\,\\% confidence limits \\citep{Wilson27}, the field detection rate in hot WDs could be anywhere between 2.8 and 30\\,\\%, while the field detection rate in cool DA WDs lies between 3.7 and 38\\,\\%. In conclusion, the data currently available are consistent with the hypothesis that weak magnetic fields occur with the same frequency in hot and cool DA WDs.  Globally, the detection of four weak magnetic fields from a total sample of 36 WDs makes it quite clear that the probability of finding a weak field in a DA WD is neither negligible, nor close to 1; at the 95\\% confidence limits, the probability lies between 4 and 25\\%.  Therefore, it appears that the probability of detecting a $\\sim 10$\\,kG field in a WD is comparable to the probability of detecting a magnetic field with strength in the range 100\\,kG -- 500\\,MG, which is $\\sim 10$\\,\\%.  Re-addressing some of the questions posed in Sect.~1, it appears now that $\\sim 10$\\,kG longitudinal fields are not ubiquitous in WDs lacking stronger fields, nor do fields seem to die away at this level. Furthermore, we have not found any significant difference between field detection rates in cool, old WDs and field detection rate in hot, young WDs. Studying these questions further will require substantially larger samples of precise field measurements than those available now.  The results of this paper highlight the need for (1) further field measurement of the MWDs already detected in this low-field regime, to fully confirm the reported detections, and to provide data on possible variabilty in order to characterise the field structures observed; (2) an extended high-precision survey of magnetic fields in hot and cool WDs, aimed at refining the frequency of occurrence of weak fields in the range studied here; and (3) a still deeper survey, using long integrations, to reach even weaker fields (note that standard errors of 300 -- 500\\,G are already achieved in a number of stars with integrations of mostly less than 30\\,min). It will also be interesting to discover whether the morphologies of the fields of kG MWDs are often roughly symmetric about the rotation axis, as seems to be the case for WD\\,2105$-$820 and as frequently happens for MWDs with stronger fields. All of these goals are within reach of observing programmes on the VLT with FORS2, although they would be very difficult on smaller telescopes.   After this paper was accepted, S. Vennes communicated to us the results of a survey of magnetic fields in a sample of 58 high proper motion white dwarfs \\citep{KawkaVennes12}. The stars of their survey are complementary to the two samples discussed in our paper. Our hot sample contains stars with typical cooling ages of 300 Myr or less, and our cool sample WDs typically have cooling ages of 300 -- 1000 Myr, while the sample of Kawka \\& Vennes is made up largely of stars with cooling ages above 1 Gyr. Because the WDs observed by Kawka \\& Vennes are both cooler and typically 2--3~mag fainter than those of our samples, their median standard error of field measurement is about 3~kG, compared to about 800~G for our sample. They are thus sensitive mainly to \\bz\\ fields larger than 10--20 kG, just above the \\bz\\ range of greatest interest to our study. However, their results seem to be significantly different from ours, as they find a probability of field detection of the order of 1 -- 2\\% per decade of field strength, while the samples discussed by us suggest probabilities of the order of 10\\% per decade in the weak-field limit. Further observations will be needed to determine if this difference is real. If the difference is indeed real, it may be an evolutionary effect of field decay with time, or a real increase in probability as we probe smaller and smaller field strengths."
    },
    "1208/1208.4246_arXiv.txt": {
        "abstract": "*{Classical Cepheids and RR Lyrae-type stars are usually considered to be textbook examples of purely radial, strictly periodic pulsators. Not all the variables, however, conform to this simple picture. In this review I discuss different forms of multi-periodicity observed in Cepheids and RR~Lyrae stars, including Blazhko effect and various types of radial and nonradial multi-mode oscillations.}  \\abstract{Classical Cepheids and RR Lyrae-type stars are usually considered to be textbook examples of purely radial, strictly periodic pulsators. Not all the variables, however, conform to this simple picture. In this review I discuss different forms of multi-periodicity observed in Cepheids and RR~Lyrae stars, including Blazhko effect and various types of radial and nonradial multi-mode oscillations.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.6243_arXiv.txt": {
        "abstract": "*{We perform a comparison between two radiative transfer algorithms commonly employed in hydrodynamical calculations of star formation: grey flux limited diffusion and the hybrid scheme, in addition we compare these algorithms to results from the Monte-Carlo radiative transfer code {\\sc mocassin}. In disc like density structures a hybrid scheme performs significantly better than the FLD method in the optically thin regions, with comparable results in optically thick regions. In the case of a forming high mass star we find the FLD method significantly underestimates the radiation pressure by a factor of $\\sim 100$. }  \\abstract{We perform a comparison between two radiative transfer algorithms commonly employed in hydrodynamical calculations of star formation: grey flux limited diffusion and the hybrid scheme, in addition we compare these algorithms to results from the Monte-Carlo radiative transfer code {\\sc mocassin}. In disc like density structures the hybrid scheme performs significantly better than the FLD method in the optically thin regions, with comparable results in optically thick regions. In the case of a forming high mass star we find the FLD method significantly underestimates the radiation pressure by a factor of $\\sim 100$. } ",
        "introduction": "\\label{sec:1} Numerical models of the star formation process have improved remarkably over the last two decades; however, many questions still remain. In particular the thermal and mechanical feedback produced by the radiation from the forming stars remains to be understood both in low mass, and in particularly  high mass star formation. Several numerical schemes for including the effects of radiation in hydrodynamic codes have been proposed including: Monte-Carlo techniques (Harries 2011); Short Characteristics (Davies et al. 2012); Flux limited diffusion and other moment methods (Levermore \\& Pomraning 1981); Pure ray-tracing techniques (Abel \\& Wandelt 2002) and Hybrid techniques - which combine various method together to arrive at a hopefully improved and faster method - (Edgar \\& Clarke 2003; Kuiper et al. 2010). While all these methods are fast enough for inclusion into a hydrodynamical calculation, the algorithm which is both fast and accurate for incorporation into a star formation calculation still remains a matter for debate. In this work I will present radiative transfer tests of the most commonly used method: Flux Limited Diffusion (FLD), and its improvement in the form of a hybrid method. ",
        "conclusions": ""
    },
    "1208/1208.4998_arXiv.txt": {
        "abstract": "By means of two- and three-dimensional particle-in-cell simulations, we investigate the process of driven magnetic reconnection at the termination shock of relativistic striped flows. In pulsar winds and in magnetar-powered relativistic jets, the flow consists of  stripes of alternating magnetic field polarity, separated by current sheets of hot plasma. At the wind termination shock, the flow compresses and the alternating fields annihilate by driven magnetic reconnection. Irrespective of the stripe wavelength $\\lambda$ or the wind magnetization $\\sigma$ (in the regime $\\sigma\\gg1$ of magnetically-dominated flows), shock-driven reconnection transfers all the magnetic energy of alternating fields to the particles, whose average Lorentz factor increases by a factor of $\\sigma$ with respect to the pre-shock value. In the limit $\\lambda/(r_L\\sigma)\\gg1$, where $r_L$ is the relativistic Larmor radius in the wind, the post-shock particle spectrum approaches a flat power-law tail with slope around $-1.5$, populated by particles accelerated by the reconnection electric field. The presence of a current-aligned ``guide'' magnetic field suppresses the acceleration of particles only when the guide field is stronger than the alternating component. Our findings place important constraints on the models of non-thermal radiation from Pulsar Wind Nebulae and relativistic jets. ",
        "introduction": "\\label{sec:intro4} Pulsar Wind Nebulae (PWNe) are bubbles of synchrotron-emitting plasma powered by the relativistic wind  of rotation-powered pulsars. Their broadband spectrum is produced by a nonthermal population of electron-positron pairs (hereafter, simply ``electrons''), presumably accelerated at the wind termination shock, where the momentum flux of the relativistic pulsar wind is balanced by the confining pressure of the nebula. The observed radio spectrum \\citeaffixed{bietenholz_97}{$F_{\\nu}\\propto\\nu^{-0.25}$ for the Crab Nebula;}  implies a distribution of emitting electrons of the form $dN/dE\\propto E^{-p}$, with a  flat slope $p\\sim1.5$. To explain the radio through optical emission of the Crab Nebula, the power law of shock-accelerated electrons should span at least three decades in particle energy \\cite{lyubarsky_03}. Flat electron spectra with slopes $p<2$ below GeV energies are also required to model the radio emission of hotspots in radio galaxies \\citep{stawarz_07} and X-ray observations of luminous blazar sources \\citep{sikora_09}.  A flat power-law spectrum with index $1<p<2$ is not  expected from the standard theory of Fermi acceleration in relativistic shocks, which normally yields slopes $p>2$ \\cite{achterberg_01,keshet_waxman_05,sironi_spitkovsky_09,sironi_spitkovsky_11a}. An acceleration mechanism capable of generating flat spectra in PWNe was discussed by \\citet{lyubarsky_03} and studied by \\citet{sironi_spitkovsky_11b}, under the assumption that the flow upstream of the termination shock consists of alternating stripes of opposite magnetic polarity, separated by current sheets of hot plasma (from now on, a ``striped wind''). For obliquely-rotating pulsars, this is the configuration expected around the equatorial plane of the wind, where the sign of the toroidal field alternates with the pulsar period. A similar geometry is invoked in the proto-magnetar model of gamma-ray bursts, if the striped structure of the equatorial wind is preserved when the flow gets redirected along the polar jet \\citep{metzger_11}. When passing through a shock (e.g., the termination shock of pulsar winds), the flow compresses and the alternating fields annihilate by shock-driven reconnection, which may produce the flat electron spectrum required by the observations.  The physics of magnetic reconnection can be captured self-consistently only by means of multi-dimensional particle-in-cell (PIC) simulations. Fully-kinetic PIC simulations provide a powerful tool to explore the microphysics of collisionless plasmas, since they can capture from first principles the fundamental interplay between charged particles and electromagnetic fields. In the context of relativistic magnetic reconnection in pair plasmas, most studies have explored the process of undriven reconnection \\citep{zenitani_01,jaroschek_04}, where field annihilation is initiated by a transient seed perturbation to an otherwise stable current sheet. As discussed above, this is not the setup expected at the  termination shock of relativistic striped winds. There, it is the shock-compression of the flow that steadily drives regions of opposite magnetic field polarity toward each other, causing reconnection. Recent experimental and numerical studies of non-relativistic plasmas have shown that driven reconnection is much faster than the undriven process \\citep{fox_11}.  In this work, we explore via multi-dimensional fully-kinetic PIC simulations the physics of driven reconnection at the termination shock of a striped relativistic electron-positron wind. We find that the alternating fields are completely dissipated upon compression by the shock, and their energy is transferred to the particles, regardless of the properties of the flow.  Broad particle spectra with  slopes $1<p<2$ are a common by-product of shock-driven reconnection, but the extent of the power-law tail depends on the wind magnetization and the stripe wavelength. ",
        "conclusions": "\\label{sec:disc} We have explored by means of 2D and 3D PIC simulations the internal structure and acceleration properties of relativistic shocks that propagate in an electron-positron striped wind, i.e., a flow consisting of stripes of alternating field polarity separated by current sheets of hot plasma.  We find that a fast MHD shock propagates into the pre-shock striped flow, compressing the incoming current sheets and initiating the process of \\textit{driven magnetic reconnection}, via the tearing-mode instability. Reconnection islands seeded by the passage of the fast shock grow and coalesce, while magnetic energy is dissipated at X-points located in between each pair of islands. When reconnection islands grow so big to occupy the entire region between neighboring current sheets, the striped structure of the flow is erased, and a hydrodynamic shock forms. By this point, the energy stored in the alternating fields has been entirely transferred to the particles via shock-driven magnetic reconnection, so the post-shock fluid behaves like an unmagnetized plasma. In our 2D and 3D simulations, we find that complete annihilation of the alternating fields occurs irrespective of the stripe wavelength $\\lambda$ or the magnetization $\\sigma$.  The main agent of particle energization is the reconnection electric field, as particles drift from a given X-point into the closest island. Whether all particles have comparable energy gains, or only a few particles are accelerated to high energies, and the majority stay cold, depends sensitively on the properties of the wind.  In the absence of guide fields, we find that the shape of the post-shock spectrum depends primarily on the combination $\\lambda/(r_L\\sigma)$, where $r_L$ is the relativistic Larmor radius in the striped wind. For small values of $\\lambda/(r_L\\sigma)$ ($\\lesssim$ a few tens), the spectrum resembles a Maxwellian distribution with mean energy $\\langle\\gamma\\rangle\\simeq\\gamma_0\\sigma$, where $\\gamma_0$ is the bulk Lorentz factor of the pre-shock flow. In the limit of very large values of $\\lambda/(r_L\\sigma)$ ($\\gtrsim$ a few hundreds), the spectrum approaches a broad power-law tail with slope $p\\simeq1.5$, extending  from  $\\gamma_{\\rm min}\\simeq\\gamma_0$ up to $\\gamma_{\\rm max}\\simeq \\gamma_0 \\sigma^{1/(2-p)}$. Our results are nearly insensitive to the strength of the guide field $B_g$, in the regime $B_g/B_0\\lesssim0.3$, where $B_0$ is the magnitude of the alternating component. For strong guide fields (i.e., $B_g/B_0\\gtrsim0.3$) the rate of particle inflow into the X-points is smaller, since the plasma should be moving across the guide field lines. This results in fewer particles being accelerated by the reconnection electric field, and the post-shock spectrum shifts to lower energies, everything else being fixed.     The particle energy spectra extracted from our simulations can be used directly to interpret the radiative signature of astrophysical nonthermal sources, and in particular of PWNe. The radio spectrum of the Crab Nebula, the prototype of PWNe, requires a population of nonthermal particles with a flat spectral slope ($p\\simeq1.5$), extending from $10^2\\unit{MeV}$ up to $10^5\\unit{MeV}$. Our results suggest that, for $p\\simeq1.5$, the particle spectrum will extend over three decades in energy if the wind magnetization at the termination shock is $\\sigma\\gtrsim30$. In addition, broad particle spectra are generated only if  $\\lambda/(r_L\\sigma)\\gtrsim$ a few tens. At the termination shock of the pulsar wind ($R=R_{TS}$) we have \\be\\label{eq:wind} \\frac{\\lambda}{r_L \\sigma}\\simeq4\\pi\\kappa \\frac{R_{LC}}{R_{TS}}~~, \\ee where $R_{LC}=c/\\Omega$ is the light cylinder radius ($\\Omega$ is the pulsar rotational frequency), and $\\kappa$ is the so-called multiplicity in the wind (i.e., the ratio of the actual density to the \\possessivecite{goldreich_julian_69} density). For the Crab, $R_{TS}\\simeq5\\times10^8R_{LC}$ \\cite{hester_02}, and most available models estimate $\\kappa\\simeq10^4-10^6$ \\citep{bucciantini_11}. Based on our findings,  the resulting value of $\\lambda/(r _L \\sigma)\\lesssim0.01$ would yield a Maxwellian-like spectrum, at odds with the wide flat spectrum required by observations. If radio-emitting electrons are produced in the equatorial plane by shock-driven reconnection, a revision of the existing estimates of $\\kappa$ is required. In this respect, we point out that the values of $\\kappa$ quoted above are averages over latitude, and one cannot exclude that particle injection into the pulsar wind is highly anisotropic, with multiplicity as large as $10^8$ along the equatorial plane. This is not in conflict with the observed lack of gradients in the radio spectral slope of the Crab \\cite{bietenholz_kronberg_92}, since strong fluid motions downstream from the termination shock would quickly fill the entire nebula with the long-lived radio-emitting electrons \\cite{delzanna_04}. Alternatively, the radio part of the spectrum may be produced at higher latitudes, where the termination shock is closer to the pulsar, and the ratio in \\eq{wind} becomes larger.  In summary, our findings place important constraints on the magnetospheric physics of pulsars, and they provide physically-grounded inputs for models of nonthermal emission in PWNe. \\ack L.S. is supported by NASA through Einstein Postdoctoral Fellowship grant number PF1-120090 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060. A.S. is supported by NSF grant AST-0807381 and NASA grant NNX10AI19G. The simulations presented in this article were performed on computational resources supported by the PICSciE-OIT High Performance Computing Center and Visualization Laboratory, and at  National Energy Research Scientific Computing Center, which is supported by the Office of Science of the US Department of Energy under contract No. DE-AC02-05CH11231.   \\appendix \\vspace{-0.1in}"
    },
    "1208/1208.4160_arXiv.txt": {
        "abstract": "We calculate the production of the gravitational waves from a double inflation model with lattice simulations. Between the two inflationary stages, gravitational waves with a characteristic frequency are produced by fluctuations of the scalar fields enhanced through parametric resonance. The wavelength of the produced gravitational waves gets extra redshift during the second inflationary stage and it can be in the observable range for the direct gravitational wave detectors. It is found that there is a possibility for the produced gravitational waves to be detected  in the planned experiments.\\\\ ",
        "introduction": "\\label{sec:intro}   The general relativity proposed by Einstein predicts the existence of the gravitational waves, which are the distortions of space-time geometry propagating through space as waves. Though it is indirect, there is strong evidence for the existence of the gravitational waves which is provided by the observation of the binary pulsar $1913+16$ \\cite{Hulse:1974eb}. Yet up to date no direct detection has been made. In the world, there are lots of ongoing and planned experiments to detect the gravitational waves directly. Laser Interferometer Gravitational-Wave Observer~(LIGO)~\\cite{LIGO} is working on the ground in USA. In Europe, Virgo~\\cite{Virgo} commenced operations in 2007 and Einstein Telescope~(ET)~\\cite{ET} is planned. In Japan, KAGRA~\\cite{KAGRA} is now under construction. It is also planned to construct the space-borne interferometers such as Laser Interferometer Space Antenna~(LISA)~\\cite{LISA},~ Deci-hertz Interferometer Gravitational Wave Observatory~(DECIGO)~ \\cite{DECIGO1,Seto:2001qf,Kawamura:2011zz}, and Big Bang Observer~(BBO)~\\cite{Crowder:2005nr}.  The direct detection experiments of the gravitational waves are expected to provide rich information about various astrophysical phenomena, but one of the ultimate goals of the experiments is to probe the early history of the universe. It is presumed that in the very early universe the vacuum energy dominates the universe and causes quasi-exponential expansion of the universe called inflation. Inflation explains the homogeneity and flatness of the universe. Planck satellite~\\cite{Ade:2013zuv} has measured anisotropies of the cosmic microwave background~(CMB) and successfully determined the cosmological parameters with high precision, which supports the prediction of the  inflationary theories. However, it is impossible to prove directly the era before photons decouple from electrons, if we use the result of the observation of cosmic microwave background only. In order to study the early universe before photon decoupling, we use the gravitational wave since the gravitational wave interacts with other particles very weakly and preserves the information about the early universe.  In this paper, we discuss the gravitational wave production from inflation. There are several primordial origins of the gravitational waves: quantum fluctuations of the metric~\\cite{Lyth:2009zz}, preheating~ \\cite{Khlebnikov:1997di,Easther:2006gt, GarciaBellido:2007af,Dufaux:2008dn}, domain walls~\\cite{Hiramatsu:2010yz,Kawasaki:2011vv}, cosmic strings~ \\cite{Berezinsky:2000vn,Kawasaki:2010yi,Damour:2000wa,Damour:2004kw}, and so on. In this paper, we focus on the preheating as the origin of gravitational waves. After the end of the inflation, the inflaton rapidly oscillates around the true vacuum, interacting with other scalar fields. At this epoch, fluctuations of the scalar fields enhanced vastly~ \\cite{GarciaBellido:1997wm,Felder:2000hj}, which produce abundant amount of gravitational waves. However, in general, the wavelength of the gravitational waves produced from reheating is so short~\\cite{Dufaux:2008dn} that we cannot detect the gravitational waves even in the future gravitational wave experiments with high sensitivity. In this work, we show that the double inflation model, which was originally proposed in~ \\cite{Randall:1995dj,GarciaBellido:1996qt, Izawa:1997df} to solve the initial value problem for the new inflation, produces gravitational waves with the wavelength long enough to be proved with the planned detectors.  In the double inflation, there are two inflationary stages. In this paper, we consider the smooth hybrid new inflation model~\\cite{Yamaguchi:2004tn}. At the first stage, smooth hybrid inflation~ \\cite{Lazarides:1995vr,Jeannerot:2000sv} occurs, followed by the second stage at which new inflation occurs. In the intermediate stage, the inflaton and waterfall fields oscillate around the minimum of the potential in the same way as the reheating. Since in the smooth hybrid inflation, the gauge symmetry of the waterfall fields is broken even during the inflation, no harmful topological defects are produced. In the previous study~\\cite{Kawasaki:2006zv}, it was found that in this model fluctuations of the scalar fields are enhanced by the parametric resonance and then primordial black holes are formed at a characteristic scale. In this paper, we pay attention to the production of the gravitational waves from these fluctuations. Since the parametric resonance is a nonlinear phenomenon, we calculate the production of the gravitational waves with lattice simulations.  The organization of this paper is as follows. We review the double inflation model in Sec.~\\ref{sec:Smooth hybrid inflation}. In Sec.~\\ref{sec:gw_production}, we calculate the amount of the produced gravitational waves by using analytical consideration and lattice simulations. We estimate the present density of the gravitational waves and the peak frequency in Sec.~\\ref{sec:present_density}. Finally, Sec.~\\ref{sec:conclusion} is devoted to the conclusion. ",
        "conclusions": "\\label{sec:conclusion}   In this paper, we have studied the production of the gravitational waves in the smooth hybrid new inflation model \\cite{Yamaguchi:2004tn,Kawasaki:2006zv}. After the smooth hybrid inflation, the inflaton and waterfall fields for the hybrid inflation start to oscillate and their fluctuations grow exponentially through the parametric resonance, which leads to the efficient production of gravitational waves. We performed the lattice simulation and calculated the amount of the gravitational waves.  From the result of the lattice simulation, we found the relation between the present gravitational energy density $\\Omega _{gw,0}$ and the peak frequency $f_{gw}$. Fig.~\\ref{fig:sensiv} shows the sensitivities of the planned detectors and gravitational wave spectrum predicted from inflation models. This figure shows that the gravitational waves produced in this model can be detectable in the future experiments in some parameter space. For example, the pink line shows the spectrum of gravitational waves produced in a set of parameters $\\mu = 3.8\\times10^{-4}M_{{\\rm pl}},~M=4.1\\times10^{-2}M_{{\\rm pl}} ,~n_s-1=-0.035$, where $n_s$ is the best-fit value of the Planck result. In this parameter set, the ultimate DECIGO can detect the gravitational waves.  \\begin{figure}[tb] \\begin{center} \\includegraphics[angle=-90,width=120mm]{./pdf/sensiv2.pdf} \\end{center} \\caption{Sensitivity curves of planned detectors, ultimate DECIGO, BBO, LISA, ET, and KAGRA. The green line shows the amplitude of the gravitational waves at the peak frequency for each set of the parameters, especially the pink line shows the spectra of the gravitational waves in the present model with $\\mu=3.8\\times10^{-4}M_{{\\rm pl}},~M=4.1\\times10^{-2}M_{{\\rm pl}}$ and $n_s-1=-0.035$. The blue line shows the spectrum of the gravitational waves from the single hybrid inflation~\\cite{Dufaux:2008dn}. The dotted line shows the upper bound on the tensor to scalar ratio. } \\label{fig:sensiv} \\end{figure}"
    },
    "1208/1208.0427.txt": {
        "abstract": "One of the most promising possibilities may be the appearance of quark matter in astrophysical phenomena in the light of recent progress in observations. The mechanism of deconfinement is not well understood, but the thermodynamical aspects of the hadron-quark (HQ) phase transition have been extensively studied in recent years. Then the mixed phase of hadron and quark matter becomes important; the proper treatment is needed to describe the HQ phase transition and derive the equation of state (EOS) for the HQ matter, based on the Gibbs conditions for phase equilibrium. We here adopt a EOS based on the baryon-baryon interactions including hyperons for the hadron phase, while we use rather simple EOS within the MIT bag model for the quark phase. %For quark matter we further try to improve the previous EOS by considering other effective models of QCD. One of the interesting consequences may be the appearance of the inhomogeneous structures called ``pasta'', which are brought about by the surface and the Coulomb interaction effects. We present here a comprehensive review of our recent works about the HQ phase transition in various astrophysical situations: cold catalyzed matter, hot matter and neutrino-trapped matter. We show how the pasta structure becomes unstable by the charge screening of the Coulomb interaction, thermal effect or the neutrino trapping effect. Such inhomogeneous  structure may affect astrophysical phenomena through its elasticity or thermal properties. Here we also discuss some implications on supernova explosion, gravitational wave and cooling of compact stars. ",
        "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %----- part 1; About QCD phase transition and phase diagram The mechanism of the hadron-quark~(HQ) deconfinement is one of hot topics in modern physics. Theoretically, numerical simulations based on the lattice gauge theory have been done to estimate the critical temperature and figure out properties of the phase transition\\cite{petreczky12}. However, at present, these attempts do not provide definite results for the deconfinement transition in high density region. % \\cite{**}. <--- Where? On the other hand, there have been done many theoretical studies by using the effective models of QCD\\cite{QCD, alf08}. Therefore, one of the most hopeful strategies to get insight into  this mechanism is to find or predict some signals about QCD phase transition in the astrophysical phenomena such as supernovae, black hole formations, mergers of neutron stars, cooling of compact stars, etc., in the light of recent theoretical and observational progress. Then the equation of state~(EOS) is a basic element to discuss the implications on above astrophysical phenomena. There have been proposed and discussed various types of scenario concerned with the HQ deconfinement transitions at high-density region, but it is still unclear, even, whether the phase transition is the cross over or the first order~\\cite{fukushima10}. Only assumption in this review is that the phase transition is of the first order as suggested by many model studies \\cite{QCD}. One of the direct consequences of this assumption is the emergence of the hadron-quark (HQ) mixed phase during the phase transition. In this review we first discuss the thermodynamical aspects of the HQ transition, and present some implications on astrophysical phenomena subsequently.  %----- part 2; First order phase transition multi component system %Maxwell Historically, the mixed phases have been often treated very naively by applying the Maxwell construction to get the EOS in thermodynamic equilibrium. %bulk Gibbs In 1990 Glendenning has pointed out that the phase equilibrium in multi-component system or in the case of plural chemical potential corresponding to the conserved quantities such as in neutron star matter must be considered based on the more general Gibbs conditions \\cite{glen, gle92}. He emphasized that the neutron-star matter should be treated as a{\\it  binary} system specified by two conserved quantities, electromagnetic charge and baryon numbers. Thus the pressure becomes no more constant as in the Maxwell construction, but still increases in the mixed phase. He also demonstrated that the mixed phase spreads in a rather wide density region within a bulk calculation, where two kinds of semi-infinite matter are considered in phase equilibrium, and the pressure balance, chemical equilibrium and the global charge-neutrality conditions are imposed, discarding the Coulomb interaction. It is important to see that particle fractions in each phase are no more the same for given average particle-fractions. Generally any phase transition in a system with two or more conserved charges must be {\\it non-congruent} \\footnote{When one dare to apply the Maxwell construction in such systems, one must violate the Gibbs conditions by ,e.g., imposing local charge neutrality.} . However, the Coulomb interaction is discarded in this calculation and we can by no means estimate how important it is, or the surface energy can not be taken into account a priori. %pasta If we consider a more realistic situation, we must consider inhomogeneous matter with various geometric structures called ``{\\it pasta}'', taking into account the finite-size effects such as the Coulomb interaction and the surface tension. The pasta resembles the {\\it nuclear pasta} in liquid-gas transition at subnuclear densities~\\cite{Lorenz1993, Oyamatsu1993, SOT1995, maru2005, new, oka}. Heiselberg et al.\\ have demonstrated that such finite-size effects are important for the mixed phase in the context of the HQ transition \\cite{HPS93}. Treating the pasta structure as in ref.~\\cite{Rav83}, they simply considered uniform particle densities in both phases. They showed a possibility of that the finite-size effects may largely restrict the region of the mixed phase. Subsequently, Voskresensky et al.\\ have shown that the rearrangement of particle densities and the screening effect for the Coulomb interaction are important, and  the large surface tension gives the mechanical instability of the geometrical structures in the pasta phase \\cite{voskre}. Using the effective model for hadron matter and the MIT bag model for quark matter, Endo et al.\\  explicitly demonstrated these features \\cite{end05}. Following this idea, we have studied the pasta structures brought by the kaon condensation and the HQ deconfinement transition \\cite{marukaon06,end05,hyp07}. More recently we have extended our framework to deal with the finite temperature case or the neutrino-trapping matter, which may be relevant to supernova explosions or neutron star mergers \\cite{yas09, yas12}.  %----- part 3; About this review In this review we first present thermodynamic concepts of the non-congruent transition of the HQ phase transition (Section \\ref{cong}). In the section we will also present some  results about the phase transition, where the effects of finite temperature or neutrino trapping are also discussed. In Section \\ref{astro} is discussed some astrophysical implications of the HQ phase transition. Section \\ref{sum} is devoted to summary and concluding remarks.     %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % \t\t\t\t\t\t\t\tSection 2. ",
        "conclusions": "\\label{sum} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  We have discussed features of the HQ mixed phase in compact stars, where multi components are in chemical equilibrium. Such phase transitions are generally not congruent and the Maxwell construction, which is a common method to obtain the equation of state (EOS) of matter with a single component, cannot be applied any more and one has to solve the Gibbs conditions instead. Using the bulk calculation, i.e., without considering the surface tension and the Coulomb interaction among two phases, we have presented basic concepts of the non-congruent transition. Although these concepts are helpful in discussing and analyzing the properties of the phase transitions in compact stars, we must bear in mind that they are obtained in an idealized and simplified situation. For realistic cases, we have to take into account rather complicated finite-size effects such as the surface tension and the Coulomb repulsion.  When the finite-size effects are present, crystalline structures of regular geometries called ``pasta structures'' emerge in the mixed phase. Charge screening, i.e., the rearrangement of charged particle densities in the presence of the Coulomb interaction is also emphasized; it sometimes causes an instability of the geometrical structures in the mixed phase. In the extreme case the EOS resembles the one given by the Maxwell construction, and uniform nuclear matter can exist in the mechanically unstable region.  %We have shown many interesting aspects of these phase transitions. %However, We note that the maximum mass of hybrid stars in our calculation is rather low in the light of a recent observation \\cite{dem10}, since contamination of hyperons considerably softens EOS. We have found that the transition of hadronic matter to quark matter suppresses the appearance of hyperons, but the maximum mass is still low due to the rather soft EOS of quark matter. Therefore, we need other EOSs of hadron and quark matter to circumvent this situation.  So far, most works on the pasta structures have used the WS cell with ansatz about the geometrical structures like droplet, rod, slab and so on. If we want to see the possibility of the intermediate shapes, or study the mixed phase in a more realistic way, Frameworks without the WS approximation or ansatz on the geometrical structures are desirable. Recently we have performed a three-dimensional calculation of non-uniform nuclear matter based on the relativistic mean-field model and the Thomas-Fermi approximation \\cite{oka}. Introducing a cubic cell with periodic boundary conditions and discretizing it into grids, we numerically solve the coupled field equations for meson mean-fields and the Coulomb potential. Randomly distributing fermions ($n, p, e$) on the grid points as an initial condition, we relax their density distributions to attain the uniformity of the chemical potentials. We have shown that typical pasta structures, obtained in the previous studies with the WS approximation, also appear in the three-dimensional calculation~\\cite{oka}.  %% %% added Possibilities of observing the signals of the mixed phase has been suggested in the spectra of the gravitational waves \\cite{SYMT2011}. We have found that one could know the existence of density discontinuity inside the star via observing the gravitational waves of not only $g$ mode but also $f$ mode. The detailed observations of $p$ mode gravitational waves could reveal the properties of the HQ mixed phase because the  $p_1$ mode gravitational waves depend on the adopted EOSs even if the density discontinuity does not exist.  QPO in giant flares  from SGR is another probe from neutron stars. If the observed QPO frequencies are associated with the torsional oscillations in neutron star crust, one might be able to obtain the information about the crust properties. In fact, we have pointed out that some of the QPO frequencies observed in giant flares might be associated with the torsional oscillations in the HQ mixed phase. %% %%  Phenomenological implications of pasta structures should deserve more elaborate studies of their thermodynamical or dynamical properties, e.g., the neutrino opacity or the thermal conductivity may be affected by the pasta structures, and their elasticity may cause a discontinuous change of internal structure of compact stars. In particular, the crystalline structure of the HQ mixed phase in the core region may bring about novel dynamical effects on compact-star phenomena by way of deformation or oscillation.    %%%%COOLING%%%%% The other important application of the EOS with the HQ mixed phase is the cooling theory of neutron stars, the long time issue discussed for decades. The cooling of neutron stars strongly depends on the neutrino emissivity of the core of stars. Simply applying the HQ mixed phase with normal quark matter, stars cool rapidly and do not match observations. However, including color superconductivity, the neutrino emissivity of the quark phase is much suppressed, and the cooling of stars with the HQ mixed phase does not conflict with observations.  Taking account of a color superconductivity in the HQ mixed phase, it became possible to have heavier stars cooling slower and lighter ones faster. We cannot get this situation without considering the HQ mixed phase. If one observes the masses of isolated neutron stars and the mass--temperature distribution can be described by the above scenario, it proves the existence of the color superconductivity in the HQ mixed phase."
    },
    "1208/1208.3430_arXiv.txt": {
        "abstract": "The measurement of muon energy is critical for many analyses in large Cherenkov detectors, particularly those that involve separating extraterrestrial neutrinos from the atmospheric neutrino background.  Muon energy has traditionally been determined by measuring the specific energy loss ($dE/dx$) along the muon's path and relating the $dE/dx$ to the muon energy.  Because high-energy muons ($E_{\\mu}~ >$~1~TeV) lose energy randomly, the spread in $dE/dx$ values is quite large, leading to a typical energy resolution of 0.29 in $\\log_{10}(E_{\\mu})$ for a muon observed over a 1~km path length in the IceCube detector.  In this paper, we present an improved method that uses a truncated mean and other techniques to determine the muon energy.  The muon track is divided into separate segments with individual $dE/dx$ values.  The elimination of segments with the highest $dE/dx$ results in an overall $dE/dx$ that is more closely correlated to the muon energy. This method results in an energy resolution of 0.22 in $\\log_{10}(E_{\\mu})$, which gives a 26\\% improvement.   This technique is applicable to any large water or ice detector and potentially to large scintillator or liquid argon detectors. ",
        "introduction": "Large ice or water Cherenkov detectors may observe up to 100,000 $\\nu_{\\mu}$ per year \\cite{Halzen:2010yj}.  These events are used for a wide variety of analyses, including searches for point sources of neutrinos \\cite{IceCube:2011ai,Abbasi:2010rd, Bogazzi:2011zza}, diffuse extraterrestrial neutrinos \\cite{Bogazzi:2011zza,Abbasi:2011jx,Dzhilkibaev:2009ja}, standard and non-standard neutrino oscillations \\cite{Abbasi:2010kx}, and measurements of the total neutrino-nucleon cross-section via neutrino absorption in the Earth.  These analyses rely upon the measurement of the neutrino energy, which is determined from the energy of the muon that is created in the neutrino interaction.  Above $\\sim$1~TeV, the muon energy is usually determined by measuring the specific energy loss, $dE/dx$, of the muon as it travels through the detector.  The Cherenkov photons from the muon and also those derived from the charged particles produced by stochastic (random) muon interactions are then detected.  This approach is disadvantageous because, for $E_{\\mu}>1$~TeV, muons lose most of their energy stochastically, and a small number of high-energy interactions will not only skew the mean but also enlarge the spread in $dE/dx$ values.  In this paper, we present an improved method for calculating the muon energy loss, which leads to significant improvement in the energy resolution.  Instead of averaging the muon $dE/dx$ over the entire observed muon path length, we divide the path into independent segments, or bins.  The $dE/dx$ is calculated separately for each bin.  Then, the bins with the highest $dE/dx$ values are discarded before calculating a new average $dE/dx$, thus producing a truncated mean.  This method is successful because the truncated mean minimizes the effects of the large stochastic events which would otherwise skew the mean and enlarge the spread.  This is the first time that the truncated mean has been systematically applied to the energy measurement of high-energy muons.  The truncated mean method has previously been calculated for muons \\cite{Auchincloss:1993zu}, although that analysis did not use it to determine the muon energy.  The method was also explored at a basic level for the DUMAND project \\cite{Mitsui:1992nt}.  The method has parallels to the one that was used to identify pions, kaons, and protons in wire chambers from their specific energy loss $dE/dx$ in the gas.  By discarding the highest 30\\% or 40\\% of the $dE/dx$ measurements from the wires, the energy resolution was greatly improved \\cite{Hauschild:1996gv,Bichsel:2006cs}.  As with wire chamber $dE/dx$ measurements, the muon energy resolution is improved by discarding the most energetic stochastic losses.  This method should be applicable to any large water or ice Cherenkov detector, such as IceCube \\cite{Halzen:2010yj}, ANTARES \\cite{Collaboration:2011nsa}, or the proposed KM3NeT \\cite{Sapienza:2011zzb}, MEMPHYS \\cite{Agostino:2012fd}, or Hyper-Kamiokande (Hyper-K) \\cite{Abe:2011ts}.  This approach may also be useful for proposed scintillator or liquid argon detectors, such as the Low Energy Neutrino Astronomy detector (LENA) \\cite{Wurm:2011zn} or the Long-Baseline Neutrino Experiment (LBNE) \\cite{Barker:2012nb}. ",
        "conclusions": "There is no apparent correlation between the energy calculation and the zenith angle (see Fig.~\\ref{fig:zenith}). The distributions of the energy residuals are very similar in shape and width.  There is also no apparent correlation between the energy calculation and azimuth angle, optical properties of the medium, or other parameters.  Thus the truncation method is universally applicable to all particle track zenith and azimuth angles within the detector, with the proper bin size. The energy resolution improves with an increasing number of bins as expected, as shown in Fig.~\\ref{fig:numbins}, but levels off at 0.18 in $\\log_{10}(E_{\\mu})$. \\begin{figure}[!h] \\centering \\includegraphics[width=0.48\\textwidth]{Fig11.eps} % \\caption{(color online) Energy resolution (in $\\log_{10}(E_{\\mu})$) for zenith angle bins of 30 degrees for $E_{\\mu}$ from 1~TeV to 100~TeV, using the truncated bins method. There is no visible correlation between zenith angle and energy resolution, for the energy range from 1~TeV to 1~EeV using an $E_{\\nu}^{-1}$ spectrum. } \\label{fig:zenith} \\end{figure} \\begin{figure}[!h] \\centering \\includegraphics[width=0.48\\textwidth]{Fig12.eps} % \\caption{(color online) Correlation between the number of bins and the average energy resolution (in $\\log_{10}(E_{\\mu})$) of the muon events, using the optimized truncated mean method, for $E_{\\mu}$ $\\scriptstyle \\gtrsim$~10~TeV. } \\label{fig:numbins} \\end{figure} \\begin{figure}[!h] \\centering \\includegraphics[width=0.48\\textwidth]{Fig13.eps} % \\caption{(color online) Comparison of the RMS and Gaussian $\\sigma$ of untruncated and truncated $dE/dx$ values vs.~$\\log_{10} (E_{\\mu}$), for simulated energies between 1~TeV and 1~EeV.  The energy resolution for the truncated method is much improved over the untruncated method.  } \\label{fig:calc_energy_res} \\end{figure}  Figure~\\ref{fig:calc_energy_res} compares the RMS and Gaussian $\\sigma$ values of truncated and untruncated $dE/dx$ for various muon energies, showing the improvement from using the truncated method.  Figures~\\ref{fig:calc_energy_orig} and \\ref{fig:calc_energy_bins} contrast the untruncated and truncated methods, respectively, for actual versus calculated muon energy.  The improvement in energy resolution is fairly uniform over the energy range.  Figure~\\ref{fig:spectrum} compares the input muon energy spectrum to the spectrum of calculated energies from the truncated method.  The agreement is quite good, with a slight shift in energy below 1~TeV due to the off-peak fit equation plus the spread in $dE/dx$ at low energies.  However, there are no glaring discontinuities in the calculated spectrum, which indicates that the truncated method gives back the original spectrum quite well.  By using the truncated mean method, the energy resolution is significantly improved. The best truncated method incorporated the following criteria: (1) only include the photoelectrons from optical modules within 10 to 80 m of the track, (2) truncate the highest 40\\% of the bins, (3) use both hit and unhit optical modules in the calculation, and (4) sum the remaining photoelectrons separately (observed and expected) to determine the new truncated $dE/dx$ value.  With these optimizations, the energy resolution was improved from 0.29 in $\\log_{10}(E_{\\mu})$ to 0.22, for the energy range of 1~TeV to 1~EeV.  This is a 26\\% improvement in the overall energy resolution, with better resolutions above 10~TeV.  The technique is applicable to any detector that uses $dE/dx$ as the primary means of the particle's energy determination.  \\begin{figure}[!h] \\centering \\includegraphics[width=0.48\\textwidth]{Fig14.eps} % \\caption{(color online) Actual muon energy versus calculated muon energy for the untruncated method with the 3-equation fit.  The black line is a perfect 1:1 correspondence. } \\label{fig:calc_energy_orig} \\end{figure} \\begin{figure}[!h] \\centering \\includegraphics[width=0.48\\textwidth]{Fig15.eps} % \\caption{(color online) Actual muon energy versus calculated muon energy for the truncated method (40\\% cuts with optimizations).  The black line is a perfect 1:1 correspondence. } \\label{fig:calc_energy_bins} \\end{figure} \\begin{figure}[!h] \\centering \\includegraphics[width=0.48\\textwidth]{Fig16.eps} % \\caption{(color online) Comparison of the truncated energy spectrum to the simulated $E_{\\nu}^{-1}$ muon energy spectrum.  The agreement is quite good between 1~TeV and 1~EeV, with a slight disagreement at low energy due to the large spread in $dE/dx$ and off-peak fit equation to convert to $E_{\\mu}$. } \\label{fig:spectrum} \\end{figure}"
    },
    "1208/1208.4026_arXiv.txt": {
        "abstract": "{The total mass of clusters of galaxies is a key parameter to study massive halos. It relates to numerous gravitational and baryonic processes at play in the framework of large scale structure formation, thus rendering its determination important but challenging.  From a sample of the 11 X-ray bright clusters selected from the \\excpres\\ sample, we investigate the optical and X-ray properties of clusters with respect to their total mass  derived from weak gravitational lensing. From multi-color wide field imaging obtained with MegaCam at CFHT, we derive the shear profile of each individual cluster of galaxies. We perform a careful investigation of all systematic sources related to the weak lensing mass determination. The weak lensing masses are then compared to the X-ray masses obtained from the analysis of \\xmm\\ observations and assuming hydrostatic equilibrium.  We find a good agreement between the two mass proxies although a few outliers with either perturbed morphology or poor quality data prevent to derive robust mass estimates. The weak lensing mass is also correlated with the optical richness and the total optical luminosity, as well as with the X-ray luminosity, to provide scaling relations within the redshift range $0.4<z<0.6$.  These relations are in good agreement with previous works at lower redshifts. For the $L_X-M$ relation we combine our sample with two other cluster and group samples from the literature, thus covering two decades in mass and X-ray luminosity, with a regular and coherent correlation between the two physical quantities.} ",
        "introduction": "Clusters of galaxies have been attracting considerable interest for their cosmological applications since a few decades. The observed cluster abundance, converted into a mass function, can be confronted to cosmological numerical simulations to put tight constraints on cosmological parameters like the amplitude of matter fluctuations or the dark energy density \\citep{bahcall92,white93,haiman01,wang04,albrecht06,mandelbaum07,rozo09}. The main limitation in the use of this mass function is the practical determination of the masses themselves.  In the simplest model of structure formation and evolution which involves a pure gravitational collapse of dark matter halos, groups and clusters are expected to form a population of self-similar objects characterized by simple relations linking the total mass to other physical quantities \\citep{kaiser86}. These scaling relations represent an efficient way to convert simple observables to masses for large samples of objects. However, the link between the total mass of an object and its baryonic tracers is not fully understood. Several processes break the self-similarity and introduce a scatter around the theoretical relation that needs to be accounted for when determining the mass function (see \\citealt{voit05} for a review). Hydrodynamic simulations offers a way to complete this model. They can define more realistic relations between the cluster total mass and the X-ray observables (e.g. \\citealt{kravtsov06,nagai07}). But the physics of the baryons within these simulations is not well constrained. Therefore, it is of prime importance to characterize and calibrate scaling relations with observational results based on accurate direct mass measurements. Such empirical relations can also improve our understanding of several non-gravitational processes disturbing the evolution of the large-scale structures.  Among the several ways to derive the total mass of a cluster of galaxies, an efficient one is to use the X-ray emission of the intra-cluster gas. The distribution in density and temperature can be measured and, assuming the hydrostatic equilibrium of the cluster, it is possible to derive the dark matter potential shape. This kind of approach has been applied successfully on many X-ray observations \\citep{pointecouteau05,vikhlinin06,buote07}.  However, masses derived from the X-ray measurements suffer systematics from non-relaxed objects or from non-thermal pressure support, and tends to underestimate the true total mass of clusters by $\\sim 10-15$\\% as seen in numerical simulations \\citep{kay04,nagai07,lau09,meneghetti10} and suggested by observational results (e.g. \\citealt{mahdavi08}).  Gravitational lensing is another way to derive the total projected mass, without any assumption about the dynamical state of the cluster. The strong lensing regime that occurs in the center of some clusters leads to elongated arcs and multiple images of background galaxies which probe the central mass distribution.  In parallel, the weak lensing regime leads to estimates of the total mass through the statistical distortion of background galaxies up to large projected distances from the cluster center (e.g. \\citealt{bartelmann01} for a review). However, this method gives reliable results only for massive clusters in which the weak lensing signal is detectable with high enough accuracy. It is also limited to intermediate redshift clusters for which the distant background galaxies are far enough behind the cluster. Gravitational lensing masses are also subject to systematics such as the determination of the redshift distribution of the sources or projection effects of matter structures near the cluster or along the line-of-sight \\citep{metzler01,hoekstra03,meneghetti10}.  Because all these mass estimators have their own weaknesses and limits and because they fail to properly recover the total mass in specific cases, there has been attempts to provide joint analysis of the mass profiles, combining strong and weak lensing \\citep{kneib03,bradac05,merten10,oguri11} or lensing and X-ray measurements \\citep{mahdavi07,morandi1423}. But the difficulty to get data of similar \"quality\" to combine them efficiently has up to now limited the use of such analysis.  In order to better quantify the scaling relations of clusters with weak lensing masses, we present in this paper the study of a sample of 11 intermediate redshift clusters which are part of the \\excpres\\ sample. These clusters have been observed in X-ray with  \\xmm\\ (Arnaud et al. in preparation) and in wide field optical imaging with Megacam at the Canada-France-Hawaii Telescope (\\citealt{foex12a}, hereafter Paper 2). The paper is organized as follows. In section 2 we present the sample of 11 clusters selected for this study. Section 3 is dedicated to the lensing methodology with special attention to the shape measurements of the background galaxies and the normalization of the lensing signal. Lensing masses are determined from the shear signal in section 4 where we also make the comparison with the X-ray masses.  Scaling relations are built and discussed in section 5 and some conclusions are drawn in section 6. In all the paper, we assume a standard $\\Lambda$-CDM cosmology with $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_{0}=70 \\ h_{70}\\ \\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$. ",
        "conclusions": "In this paper we provide for the first time a thorough weak lensing analysis of a sample of 11 galaxy clusters (a sub-sample of the \\excpres\\ sample -- Arnaud et al. in preparation) in a relatively high redshift range, $0.4<z<0.6$.  We have conducted a careful validation of our weak lensing pipeline, thanks in particular to the STEP simulations which allowed to re-calibrate our shear measurements. The shear profiles were fitted with a NFW profile with a fixed concentration parameter $c_{200}=4$. This is due to the lack of good constraints on the shear profile near the center so only the total mass $M_{200}$ is a reliable measure.  These weak lensing masses were compared to the X-ray masses derived from  \\xmm\\ data. Over the 11 clusters, we find a good match between these two mass estimators for 7 clusters while we found 4 outliers, thus having almost a $1\\sigma$ agreement in the two methods. Such results are consistent with other similar studies performed for lower redshift samples.  The primary goal of this study was to look for correlations between the total mass of galaxy clusters and several baryonic tracers. In our sample of massive clusters, we found that the total mass is proportional to the optical richness of galaxies.  This result is in good agreement with previous works at lower redshifts \\citep{becker07,johnston07,reyes08,rykoff08b,mandelbaum08,rozo09}. We reach the same conclusion for the correlation between mass and optical luminosity, with a constant ratio $M/L$, even if the correlation is weaker with a larger intrinsic dispersion.  Using \\xmm\\ data and the weak lensing masses, we build scaling relations with the X-ray total luminosity and the global spectroscopic temperature. Both correlations, although weakly constrained due to limited coverage of the mass range, give results in good agreement with previous studies: (i) the mass-luminosity relation presents a non self-similar slope, larger than expected from a purely gravitational model of structure formation. (ii) the mass-temperature relation is roughly compatible with the hierarchical prediction. We have extended our investigation of the $L_X-M$ relation by combining our cluster sample with two samples of groups and clusters for which weak lensing masses were obtained \\citep{rykoff08b,leauthaud10}. The resulting $L_X-M_{200}$ relation covers nearly two decades in mass and displays a regular shape with a well constrained slope.  Even if the limited size of our sample and its reduced mass range coverage allow us to only probe the high end of the mass distribution of clusters, we have demonstrated the synergy with the X-ray measurements (i.e., mass, luminosity and temperature). More specifically, we have confirmed that the scaling relations with the lensing mass proxy are already in place at $z \\sim 0.5$, with, at first order, no significant departures from the same relations at lower redshift (assuming a self-similar evolution). The limited quality of the lensing data for our sample restricts the added value of a joint analysis, but the weak lensing and X-ray coverage of cluster prove to be complementary. Indeed X-ray data allow, to efficiently probe the central regions out to about $R_{500}$. Beyond this radius the X-ray brightness rapidly drops, whereas the weak-lensing signal picks up to characterize the outer parts of massive halos. A perspective work would be to investigate the strong lensing signal (see Table\\ref{table:clusters}) of each cluster in our sample, in order to perform a full lensing analysis from the inner part (with the strong lensing signal) to the outer parts (with the weak lensing signal). This would provide a coherent mass proxy estimator to be compared directly with X-ray and dynamical proxies.  Finally, in the perspective of large optical surveys such as the Euclid space mission, self calibrated mass proxies will be needed for the full scientific exploitation of the mission. The $M-N$ is an obvious candidate, which tight calibration will require validation against other relations and mass proxies, such as the X-ray mass. The $M-N$ relation we provide in this paper is an early result in the redshift range $0.4<z<0.6$. It illustrates the difficulty to constrain this particular scaling relation. Provided the $N$ estimator is made as universal as possible, which can be a non trivial task, this relation is promising in the framework of optical surveys."
    },
    "1208/1208.0369.txt": {
        "abstract": "{We present a sample of 209 variable objects -- very likely optical counterparts to the X-ray sources detected in the direction of the Galactic center by the Galactic Bulge Survey (GBS) carried out with the Chandra satellite. The variable sources were found in the databases of the OGLE long term survey monitoring regularly the Galactic bulge since 1992. The counterpart candidates were searched based on the X-ray source position in the radius of 3\\zdot\\arcs9.  Optical light curves of the candidates comprise a full variety of variability types: spotted stars, pulsating red giants (potentially secondary stars of symbiotic variables), cataclysmic variables, eclipsing binary systems, irregular non-periodic objects including an AGN (GRS 1734--292).  Additionally, we find that positions of 19 non-variable stars brighter than 16.5~mag in the OGLE databases are so well aligned with the X-ray positions ($<0\\zdot\\arcs75$) that these objects are also likely optical counterparts to the GBS X-ray sources.  We provide the OGLE astrometric and photometric information for all selected objects and their preliminary classifications. Photometry of the candidates is available from the OGLE Internet archive.}{X-ray: stars -- X-ray: binaries -- Stars: activity -- binaries: symbiotic -- novae, cataclysmic variables -- Galaxies: Seyfert} ",
        "introduction": "Recent years have been witnessing an enormous progress in the X-ray astronomy thanks to many satellite missions with continuously more sensitive state-of-the-art detectors. The number of detected X-ray sources increased rapidly. For example, the surveys of the Magellanic Clouds led to the discovery of tens of new X-ray pulsars (Coe \\etal 2005) and AGNs (Koz\u00b3owski \\etal 2012), Galactic center survey discovered more than thousand new X-ray sources toward the Galactic center (Jonker \\etal 2011), the XMM-Newton catalog of serendipitous X-ray sources includes almost 200\\,000 objects (Watson \\etal 2009),  Proper classification and then interpretation of an X-ray source usually requires additional observations in other wavelength bands, usually in the optical or IR range. Therefore, it is crucial to find the X-ray source counterparts in these bands, if the data exist.  The accuracy of positions of X-ray sources detected by the modern instruments is comparable to that obtained in the optical range. This in principle can make the cross-identification of the optical counterparts relatively straightforward. Indeed, it is often true in empty stellar fields, where the X-ray position unambiguously points to an object. The situation is, however, much more difficult in dense stellar fields such as the Magellanic Clouds or Galactic center, where one can detect even several dozen objects within the X-ray position error box. In such cases positive cross-identification can often be non-trivial.  On the other hand, the photometric variability of an object located close to the X-ray source position can be a very good indicator of its close relation with the detected X-ray source. It is well known that many different types of variable objects emit X-rays, for example, all types of cataclysmic variables, low and high mass X-ray binaries, chromospherically active and spotted stars, long period variables etc. Thus, the detection of optical variability in an object located in the X-ray position error box practically proves correctness of the cross-identification.  What is more important, variability provides many additional information on the X-ray system for its proper classification and detailed studies of its structure. For example, the detected periodicities can be interpreted as related to orbital or rotational periods. Observed flares or flickering can provide information on the rapid phenomena occurring in accreting systems or on stellar chromospherical activity.  Here, we present the results of our search for variable objects located at the positions of the X-ray sources detected by the Galactic Bulge Survey (GBS, Jonker \\etal 2011). Altogether, we have found 209 variable optical objects that can be considered as the optical counterparts to X-ray sources detected by the GBS survey. For all of them we provide their basic classification and light curve parameters. The photometric data are available from the OGLE Internet archive (see Section~5). ",
        "conclusions": "\\vspace*{5pt} We performed tentative classification of each discovered variable star based on its photometric behavior. Initially, the variables were divided into two broad categories: those with periodic and irregular variability. From the former group a few sub-categories, sometimes overlapping, were then assigned.  The most numerous sub-group -- 81 objects -- consists of spotted stars with characteristic pattern of photometric behavior resembling RS~CVn type stars. In most cases, we register only the variability caused by spots (rotation, almost synchronized with orbital period in the binary systems) and their slow migration on the stellar surface. However, the classical eclipsing RS~CVn type stars are also present in our sample. Typical light curves of spotted stars are shown in Fig.~2. \\begin{figure}[p] \\hglue-7mm{\\includegraphics[width=13cm, bb=100 125 505 710]{fig2.ps}} \\FigCap{Light curves of spotted stars. Two {\\it upper panels} show time series and phased with the period 12.771~days light curve of the optical counterpart candidate (BLG500.27 22801) to the GBS \\#29 X-ray source. Two {\\it bottom panels} show the light curves of the eclipsing system BLG501.29 138642 ($P=11.9936$~days) -- a counterpart to the GBS \\#701 X-ray source.} \\end{figure}  The second group consists of pulsating red giants (OGLE Small Amplitude Red Giants, OSARGs, SemiRegular Variables, SRVs, Miras - \\cf Soszy\u00f1ski \\etal 2011b) -- 25 objects. Such stars can be members of X-ray binary systems with a compact companion and accreting processes. Fig.~3 presents typical light curves of objects from this sub-category.  \\begin{figure}[p] \\hglue-7mm{\\includegraphics[width=13cm, bb=100 125 505 710]{fig3.ps}} \\FigCap{Light curves of pulsating red giants. Two {\\it upper panels} show time series and phased with the period 42.735 days light curve of the OSARG-type optical counterpart candidate (BLG501.24 18825) to the GBS \\#715 X-ray source. Two {\\it bottom panels} show the light curves of the SRV candidate (BLG504.23 59042; $P=172.4$~days) to the GBS \\#907 X-ray source.} \\end{figure}  Eclipsing stars subgroup includes 36 objects with clear eclipses. Contact and detached short period binary systems are usually chromospherically active due to fast rotation tidally synchronized with the short orbital period and as a consequence -- strong X-ray sources. Upper panels of Fig.~4 show the sample light curves of stars from this class. \\begin{figure}[p] \\hglue-7mm{\\includegraphics[width=13cm, bb=100 125 505 710]{fig4.ps}} \\FigCap{Phased light curves of eclipsing and periodic counterpart candidates to the GBS X-ray sources: GBS \\#684 (BLG500.09 136274), GBS \\#152 (BLG501.25 43802), GBS \\#289 (BLG504.16 90241) and GBS \\#428 (BLG500.11 129634).} \\end{figure} \\begin{figure}[htb] \\includegraphics[width=11.3cm, bb=100 280 505 710]{fig5.ps} \\FigCap{Time series of the irregular counterpart candidates to the GBS X-ray sources: GBS \\#2 (BLG653.01 89053, Seyfert galaxy GRS 1734--292), GBS \\#212 (BLG654.22.17003) and GBS \\#332 (BLG654.20.36111).} \\end{figure} \\begin{figure}[htb] \\includegraphics[width=11.5cm, bb=100 280 505 710]{fig6.ps} \\FigCap{Light curves of cataclysmic variables -- optical counterpart candidates to the X-ray sources: GBS \\#18 (BLG675.14 49163), GBS \\#714 (BLG501.15 11770) and GBS \\#426 (BLG501.28 111658).} \\end{figure} \\begin{figure}[htb] \\includegraphics[width=11.3cm, bb=100 280 505 710]{fig7.ps} \\FigCap{Light curves of optical counterpart candidates to the X-ray source GBS \\#19. {\\it Upper} and {\\it middle panels:} eclipsing system BLG501.24 34934. {\\it Lower panel:} pulsating red giant BLG501.24 34423.} \\end{figure}  Two RR~Lyr pulsating stars, one of RRab and one of RRc type, were also found in the selected sample of variable stars. Although pulsating stars (Cepheids, RR~Lyr) can potentially be soft X-ray emitters only a handful of Cepheids and RR Lyr stars have been detected in the X-ray band so far (Engle and Guinan 2012). Our detection of the coincidence of the positions of two RR~Lyr stars with X-ray sources, could be thus an important discovery. However, we suspect that this may be a by chance coincidence. The number of RR~Lyr stars detected in the typical Galactic bulge field analyzed in this paper is about 40 per $2048\\times4102$ pixel OGLE-IV CCD chip according to Soszy\u00f1ski \\etal (2011a) OGLE Catalog of Variable Stars. Thus, we may expect about two RR~Lyr variable star detections in the probed circle of 15 pixel radius in about 700 random directions toward the Galactic center. Close examination of the finding charts also suggest a by chance coincidence -- both RR~Lyr stars are relatively far from the X-ray position, one almost at the boundary of the searched circle.  The variability of the remaining periodic stars (48 objects) may be caused either by the orbital motion (ellipsoidal variations), if they are binaries, or by the rotation. In the case of ellipsoidal variation the real period is twice of that found. Typical examples of stars from this group are shown in the bottom panels of Fig.~4.  Irregular variables form a separate group. This kind of variability may be an indicator of accreting systems like high or low mass X-ray binaries or nova-like variables (flickering). It cannot be ruled out, however, that some of the objects from this subgroup are actually quasi-periodic but the small number of observations collected so far prevented sound period detection.  It is also worth noticing that the optical counterpart to the GBS \\#2 source, which is a known Seyfert galaxy, GRS 1734--292 (Mart{\\'\\i} \\etal 1998), is classified as irregular variable. OGLE provides, thus, long term optical monitoring of this interesting AGN. Three examples of irregular variables from this subgroup, including the counterpart of GBS \\#2, are shown in Fig.~5.  We also found three classical cataclysmic variables, two typical dwarf novae and one Z~Cam type star. Cataclysmic variables are well known X-ray emitters so the detection of these stars near the X-ray sources position clearly indicates that they are optical counterparts to these X-ray sources. Fig.~6 shows the light curves of these eruptive variables.  Table~1 lists all the X-ray sources from the GBS catalog cross-identified with OGLE variable objects. For each optical counterpart candidate the OGLE equatorial coordinates and distance from the X-ray source (in arcsecs) are provided, its mean optical photometry ({\\it I}-band magnitude, $V-I$ color, if available), period in days in the case of periodic variable stars and classification. Entries of counterpart candidates found in the OGLE-II databases are marked with italic font.  \\renewcommand{\\TableFont}{\\scriptsize} \\MakeTableSep{r@{\\hspace{7pt}}c@{\\hspace{3pt}}r@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{5pt}}r@{\\hspace{0pt}}c@{\\hspace{0pt}}l@{\\hspace{5pt}}l} {12.5cm}{Optical counterparts to the X-ray sources detected by the Galactic Bulge Survey revealing optical variability} {\\hline \\noalign{\\vskip3pt} GBS               & OGLE-IV & OGLE-IV              &    RA    &   DEC    & $D$     & $I$   & $V-I$ & \\multicolumn{3}{c}{$P$}    & Remarks\\\\ \\multicolumn{1}{c}{ID} & Field   &\\multicolumn{1}{c}{No}& [2000.0] & [2000.0] & [\\arcs] & [mag] & [mag] & \\multicolumn{3}{c}{[days]} &\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} 2 & BLG653.01 &  89053 & 17\\uph37\\upm28\\zdot\\ups38 & $-29\\arcd08\\arcm02\\zdot\\arcs1$ & 0.16 & 16.422 & --   &   --&&       & I AGN\\\\ 16 & BLG504.24 &     65 & 17\\uph55\\upm45\\zdot\\ups84 & $-27\\arcd58\\arcm13\\zdot\\arcs9$ & 0.18 & 14.572 & 2.383 &   4&.&95     & SP\\\\ 18 & BLG675.14 &  49163 & 17\\uph39\\upm35\\zdot\\ups76 & $-27\\arcd29\\arcm35\\zdot\\arcs9$ & 0.15 & 17.819 & --   &   --&&       & CV\\\\ {\\bf  19} & BLG501.24 &  34423 & 17\\uph49\\upm54\\zdot\\ups52 & $-29\\arcd43\\arcm35\\zdot\\arcs8$ & 0.66 & 15.668 & 5.938 &  31&.&65     & OSARG\\\\ {\\bf  19} & BLG501.24 &  34934 & 17\\uph49\\upm54\\zdot\\ups63 & $-29\\arcd43\\arcm35\\zdot\\arcs4$ & 0.87 & 17.983 & 1.519 &   0&.&3587   & E\\\\ 23 & BLG675.02 &  36509 & 17\\uph42\\upm31\\zdot\\ups57 & $-27\\arcd43\\arcm48\\zdot\\arcs3$ & 0.88 & 14.960 & --   &  44&.&63     & OSARG\\\\ 24 & BLG501.16 &  59134 & 17\\uph48\\upm49\\zdot\\ups50 & $-30\\arcd01\\arcm10\\zdot\\arcs3$ & 0.40 & 12.342 & 1.711 &  24&.&038    & SP\\\\ 29 & BLG500.27 &  22801 & 17\\uph53\\upm41\\zdot\\ups90 & $-28\\arcd03\\arcm53\\zdot\\arcs7$ & 0.37 & 15.619 & 2.134 &  12&.&771    & SP\\\\ 56 & BLG501.32 &   6608 & 17\\uph50\\upm00\\zdot\\ups36 & $-29\\arcd24\\arcm12\\zdot\\arcs7$ & 1.23 & 13.706 & 2.558 &  21&.&436    & E SP\\\\ 65 & BLG500.13 &  14897 & 17\\uph51\\upm35\\zdot\\ups26 & $-28\\arcd43\\arcm45\\zdot\\arcs3$ & 0.42 & 12.977 & 0.909 &   0&.&70676  & SP\\\\ 71 & BLG675.06 &   5524 & 17\\uph39\\upm49\\zdot\\ups29 & $-27\\arcd54\\arcm20\\zdot\\arcs6$ & 0.28 & 14.056 & --   &  --&&        & I\\\\ 89 & BLG653.11 &  44462 & 17\\uph36\\upm24\\zdot\\ups28 & $-28\\arcd45\\arcm32\\zdot\\arcs7$ & 0.24 & 14.228 & --   &  18&.&382    & SP\\\\ 103 & BLG501.29 &  50854 & 17\\uph51\\upm55\\zdot\\ups79 & $-29\\arcd23\\arcm10\\zdot\\arcs1$ & 3.72 & 14.714 & 1.684 &  --&&        & I\\\\ 106 & BLG500.08 & 191736 & 17\\uph54\\upm34\\zdot\\ups47 & $-28\\arcd41\\arcm48\\zdot\\arcs9$ & 0.37 & 12.908 & 1.676 &  32&.&57329  & SP\\\\ 122 & BLG653.08 &  10840 & 17\\uph38\\upm37\\zdot\\ups16 & $-28\\arcd45\\arcm44\\zdot\\arcs9$ & 0.44 & 13.369 & --   &  43&.&668    & SP\\\\ {\\it   124}&{\\it BUL\\_SC5}& {\\it 134677} & {\\it 17\\uph50\\upm07\\zdot\\ups00} & ${\\it -30\\arcd01\\arcm53\\zdot\\arcs9}$ & {\\it 0.77} & {\\it 12.481} & {\\it 2.097} &  --&&  & {\\it I}\\\\ 127 & BLG500.18 &  12184 & 17\\uph54\\upm21\\zdot\\ups19 & $-28\\arcd30\\arcm02\\zdot\\arcs2$ & 2.87 & 16.552 & 3.901 & 144&.&93     & P\\\\ 133 & BLG501.29 &  45257 & 17\\uph52\\upm01\\zdot\\ups41 & $-29\\arcd25\\arcm07\\zdot\\arcs9$ & 1.04 & 13.191 & 3.700 &  19&.&292    & OSARG\\\\ 134 & BLG648.28 &  31469 & 17\\uph47\\upm14\\zdot\\ups97 & $-26\\arcd16\\arcm49\\zdot\\arcs0$ & 0.05 & 13.711 & --   &  18&.&526    & OSARG\\\\ 137 & BLG504.14 & 157126 & 17\\uph55\\upm53\\zdot\\ups26 & $-28\\arcd16\\arcm33\\zdot\\arcs9$ & 0.69 & 15.113 & 1.315 &   0&.&43104  & E SP\\\\ 143 & BLG675.26 &  84886 & 17\\uph42\\upm49\\zdot\\ups73 & $-26\\arcd48\\arcm23\\zdot\\arcs0$ & 0.23 & 13.946 & --   &   1&.&72473  & P\\\\ 144 & BLG654.25 &  20335 & 17\\uph33\\upm00\\zdot\\ups32 & $-29\\arcd36\\arcm14\\zdot\\arcs9$ & 0.47 & 14.469 & --   &   2&.&47771  & E\\\\ 152 & BLG501.25 &  43802 & 17\\uph48\\upm56\\zdot\\ups72 & $-29\\arcd38\\arcm35\\zdot\\arcs9$ & 1.05 & 15.461 & 1.384 &   0&.&378895 & E\\\\ 155 & BLG504.07 &  23237 & 17\\uph55\\upm42\\zdot\\ups07 & $-28\\arcd27\\arcm39\\zdot\\arcs5$ & 0.69 & 16.515 & 2.057 &   0&.&53454  & E\\\\ 163 & BLG504.15 &  66246 & 17\\uph55\\upm32\\zdot\\ups43 & $-28\\arcd04\\arcm31\\zdot\\arcs4$ & 0.33 & 15.041 & 2.501 &   2&.&02896  & E\\\\ 174 & BLG654.27 &  54219 & 17\\uph36\\upm58\\zdot\\ups35 & $-29\\arcd23\\arcm51\\zdot\\arcs7$ & 1.33 & 12.848 & --   & 217&.&4      & P\\\\ 180 & BLG500.20 & 154077 & 17\\uph52\\upm20\\zdot\\ups78 & $-28\\arcd23\\arcm40\\zdot\\arcs6$ & 0.56 & 13.544 & --   &   2&.&5896   & E\\\\ 185 & BLG654.21 &  59612 & 17\\uph35\\upm27\\zdot\\ups56 & $-29\\arcd45\\arcm16\\zdot\\arcs6$ & 0.53 & 12.936 & --   &  44&.&25     & P\\\\ 190 & BLG653.12 &      5 & 17\\uph35\\upm44\\zdot\\ups39 & $-28\\arcd54\\arcm00\\zdot\\arcs4$ & 0.92 & 14.678 & --   & 225&.&2      & MIRA\\\\ 195 & BLG654.30 &  14819 & 17\\uph35\\upm07\\zdot\\ups54 & $-29\\arcd19\\arcm05\\zdot\\arcs1$ & 0.29 & 14.712 & --   &  28&.&5      & E\\\\ 196 & BLG500.07 &   1605 & 17\\uph50\\upm03\\zdot\\ups35 & $-29\\arcd11\\arcm50\\zdot\\arcs4$ & 1.28 & 17.089 & 1.854 &   1&.&9972   & SP\\\\ 200 & BLG504.16 &  58311 & 17\\uph54\\upm45\\zdot\\ups46 & $-28\\arcd07\\arcm43\\zdot\\arcs6$ & 2.79 & 18.350 & --   &  72&.&464    & P\\\\ 203 & BLG653.05 &  39870 & 17\\uph34\\upm55\\zdot\\ups60 & $-28\\arcd58\\arcm56\\zdot\\arcs9$ & 2.11 & 13.949 & --   &   9&.&38086  & SP\\\\ 204 & BLG500.10 & 172048 & 17\\uph53\\upm12\\zdot\\ups45 & $-28\\arcd47\\arcm46\\zdot\\arcs5$ & 0.97 & 15.267 & 1.692 &   2&.&014    & SP\\\\ 212 & BLG654.22 &  17003 & 17\\uph35\\upm10\\zdot\\ups15 & $-29\\arcd36\\arcm06\\zdot\\arcs1$ & 0.78 & 15.209 & --   &  --&&        & I\\\\ 216 & BLG653.09 &  79513 & 17\\uph37\\upm21\\zdot\\ups71 & $-28\\arcd53\\arcm57\\zdot\\arcs3$ & 0.32 & 13.216 & --   &  16&.&42036  & SP\\\\ 222 & BLG648.16 &   1012 & 17\\uph44\\upm05\\zdot\\ups45 & $-27\\arcd00\\arcm13\\zdot\\arcs3$ & 0.11 & 15.514 & --   &  46&.&083    & SP\\\\ 227 & BLG501.28 & 155033 & 17\\uph52\\upm19\\zdot\\ups58 & $-29\\arcd21\\arcm09\\zdot\\arcs8$ & 0.08 & 14.191 & 2.300 &   2&.&947    & SP\\\\ 228 & BLG534.32 &  97586 & 17\\uph49\\upm24\\zdot\\ups63 & $-30\\arcd34\\arcm44\\zdot\\arcs7$ & 0.42 & 13.089 & 1.835 &  29&.&916    & SP\\\\ 231 & BLG648.28 &  36747 & 17\\uph47\\upm14\\zdot\\ups73 & $-26\\arcd09\\arcm38\\zdot\\arcs2$ & 0.25 & 13.140 & --   &  69&.&32     & OSARG\\\\ 234 & BLG675.28 &  21714 & 17\\uph41\\upm47\\zdot\\ups68 & $-26\\arcd46\\arcm49\\zdot\\arcs8$ & 2.86 & 16.539 & --   &  26&.&455    & P\\\\ 237 & BLG654.17 &  56669 & 17\\uph38\\upm15\\zdot\\ups98 & $-29\\arcd39\\arcm39\\zdot\\arcs9$ & 0.72 & 14.138 & --   &   0&.&273483 & E\\\\ 240 & BLG654.21 &   7912 & 17\\uph35\\upm47\\zdot\\ups78 & $-29\\arcd43\\arcm33\\zdot\\arcs2$ & 0.56 & 12.523 & --   &  --&&        & I\\\\ 253 & BLG504.25 &  87539 & 17\\uph54\\upm27\\zdot\\ups60 & $-27\\arcd47\\arcm32\\zdot\\arcs9$ & 0.89 & 12.689 & 0.943 &   1&.&31734  & E SP\\\\ 282 & BLG653.11 &    550 & 17\\uph36\\upm34\\zdot\\ups43 & $-28\\arcd53\\arcm33\\zdot\\arcs9$ & 0.05 & 18.322 & --   &   9&.&72760  & P\\\\ {\\bf 289} & BLG504.16 &  90151 & 17\\uph54\\upm39\\zdot\\ups15 & $-28\\arcd08\\arcm50\\zdot\\arcs0$ & 0.32 & 14.723 & 2.400 &  11&.&84     & SP\\\\ {\\bf 289} & BLG504.16 &  90241 & 17\\uph54\\upm38\\zdot\\ups88 & $-28\\arcd08\\arcm51\\zdot\\arcs4$ & 3.56 & 15.575 & 4.250 &  25&.&38     & P\\\\ 295 & BLG500.10 & 186047 & 17\\uph53\\upm11\\zdot\\ups01 & $-28\\arcd42\\arcm27\\zdot\\arcs7$ & 0.25 & 13.432 & 1.689 &   2&.&92227  & SP\\\\ \\noalign{\\vskip3pt} \\hline} \\setcounter{table}{0} \\MakeTableSepp{r@{\\hspace{7pt}}c@{\\hspace{3pt}}r@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{5pt}}r@{\\hspace{0pt}}c@{\\hspace{0pt}}l@{\\hspace{5pt}}l} {12.5cm}{Continued} {\\hline \\noalign{\\vskip3pt} GBS               & OGLE-IV & OGLE-IV              &    RA    &   DEC    & $D$     & $I$   & $V-I$ & \\multicolumn{3}{c}{$P$}    & Remarks\\\\ \\multicolumn{1}{c}{ID} & Field   &\\multicolumn{1}{c}{No}& [2000.0] & [2000.0] & [\\arcs] & [mag] & [mag] & \\multicolumn{3}{c}{[days]} &\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} 300 & BLG654.14 &  66787 & 17\\uph33\\upm51\\zdot\\ups57 & $-30\\arcd06\\arcm30\\zdot\\arcs9$ & 0.62 & 15.282 & --   &  21&.&1417   & SP\\\\ 303 & BLG504.31 &  31226 & 17\\uph56\\upm02\\zdot\\ups11 & $-27\\arcd39\\arcm21\\zdot\\arcs2$ & 2.17 & 12.753 & 1.213 &   0&.&8313   & SP\\\\ 304 & BLG504.31 &  31226 & 17\\uph56\\upm02\\zdot\\ups11 & $-27\\arcd39\\arcm21\\zdot\\arcs2$ & 1.57 & 12.753 & 1.213 &   0&.&8313   & SP\\\\ 326 & BLG675.20 &  61377 & 17\\uph41\\upm31\\zdot\\ups36 & $-27\\arcd09\\arcm59\\zdot\\arcs2$ & 0.32 & 13.827 & --   &   0&.&23462  & P\\\\ 330 & BLG653.19 &  81200 & 17\\uph36\\upm43\\zdot\\ups88 & $-28\\arcd21\\arcm21\\zdot\\arcs3$ & 0.82 & 15.069 & --   &  --&&        & I\\\\ 331 & BLG654.28 &  28452 & 17\\uph36\\upm23\\zdot\\ups47 & $-29\\arcd22\\arcm34\\zdot\\arcs7$ & 0.14 & 17.277 & --   &  18&.&215    & SP\\\\ 332 & BLG654.20 &  36111 & 17\\uph36\\upm20\\zdot\\ups20 & $-29\\arcd33\\arcm38\\zdot\\arcs9$ & 0.80 & 15.156 & --   &  --&&        & I\\\\ 336 & BLG654.04 &  45052 & 17\\uph35\\upm31\\zdot\\ups42 & $-30\\arcd24\\arcm23\\zdot\\arcs3$ & 0.38 & 15.696 & --   &   9&.&99     & SP\\\\ 354 & BLG500.18 & 172616 & 17\\uph53\\upm52\\zdot\\ups75 & $-28\\arcd25\\arcm13\\zdot\\arcs2$ & 0.79 & 15.092 & 2.074 &   8&.&21018  & SP\\\\ {\\it   362}& {\\it BUL\\_SC5}&{\\it 284655} & {\\it 17\\uph50\\upm37\\zdot\\ups53} & ${\\it -29\\arcd40\\arcm45\\zdot\\arcs5}$ & {\\it 0.26} & {\\it 11.420} & {\\it 2.272} & {\\it 107}&.&{\\it 53} & {\\it P}\\\\ 363 & BLG501.31 &  34121 & 17\\uph50\\upm33\\zdot\\ups73 & $-29\\arcd21\\arcm33\\zdot\\arcs5$ & 0.35 & 14.005 & 1.679 &   8&.&7413   & SP\\\\ 375 & BLG675.17 &  11128 & 17\\uph43\\upm57\\zdot\\ups95 & $-27\\arcd10\\arcm34\\zdot\\arcs0$ & 1.00 & 15.365 & --   &  62&.&8      & SP\\\\ 387 & BLG654.27 &  27852 & 17\\uph37\\upm05\\zdot\\ups58 & $-29\\arcd24\\arcm52\\zdot\\arcs5$ & 0.86 & 13.044 & --   &   4&.&831    & P\\\\ 391 & BLG654.11 &  47010 & 17\\uph36\\upm06\\zdot\\ups42 & $-30\\arcd03\\arcm11\\zdot\\arcs0$ & 0.33 & 13.977 & --   &  11&.&47     & E SP\\\\ 395 & BLG654.31 &  30300 & 17\\uph34\\upm18\\zdot\\ups96 & $-29\\arcd28\\arcm33\\zdot\\arcs8$ & 1.12 & 14.868 & --   &   5&.&965    & SP\\\\ 397 & BLG654.23 &  61348 & 17\\uph33\\upm59\\zdot\\ups54 & $-29\\arcd49\\arcm47\\zdot\\arcs2$ & 0.50 & 13.190 & --   &   0&.&77471  & E\\\\ 414 & BLG504.30 &   6878 & 17\\uph57\\upm09\\zdot\\ups75 & $-27\\arcd32\\arcm49\\zdot\\arcs7$ & 0.25 & 13.715 & 1.585 &   0&.&40826  & SP\\\\ 416 & BLG504.16 &  83867 & 17\\uph54\\upm35\\zdot\\ups04 & $-28\\arcd11\\arcm52\\zdot\\arcs5$ & 1.43 & 14.998 & --   &  16&.&5563   & P\\\\ 417 & BLG504.16 & 130997 & 17\\uph54\\upm27\\zdot\\ups75 & $-28\\arcd07\\arcm49\\zdot\\arcs3$ & 0.29 & 13.749 & 0.532 &   0&.&95923  & P\\\\ 426 & BLG501.28 & 111658 & 17\\uph52\\upm36\\zdot\\ups05 & $-29\\arcd19\\arcm39\\zdot\\arcs9$ & 0.17 & 17.293 & 0.439 &  --&&        & CV\\\\ 428 & BLG500.11 & 129634 & 17\\uph52\\upm32\\zdot\\ups88 & $-28\\arcd37\\arcm48\\zdot\\arcs5$ & 0.55 & 17.698 & 1.710 &   2&.&56279  & P\\\\ 429 & BLG500.20 & 120943 & 17\\uph52\\upm32\\zdot\\ups43 & $-28\\arcd21\\arcm01\\zdot\\arcs1$ & 0.28 & 16.839 & 2.037 &  10&.&94092  & SP\\\\ 432 & BLG500.13 &  29005 & 17\\uph51\\upm19\\zdot\\ups17 & $-28\\arcd53\\arcm08\\zdot\\arcs9$ & 0.49 & 14.163 & 1.781 &  11&.&876    & SP\\\\ 434 & BLG500.05 &  84278 & 17\\uph51\\upm10\\zdot\\ups63 & $-29\\arcd01\\arcm17\\zdot\\arcs7$ & 0.62 & 14.479 & 1.101 &   0&.&69871  & SP\\\\ 439 & BLG534.25 &  65908 & 17\\uph48\\upm40\\zdot\\ups65 & $-30\\arcd55\\arcm53\\zdot\\arcs2$ & 1.75 & 16.443 & --   &  64&.&516    & SRV\\\\ 472 & BLG675.22 &  60741 & 17\\uph40\\upm25\\zdot\\ups65 & $-27\\arcd05\\arcm40\\zdot\\arcs4$ & 0.29 & 14.623 & --   &  21&.&834    & SP\\\\ 483 & BLG654.26 &  59596 & 17\\uph37\\upm31\\zdot\\ups10 & $-29\\arcd24\\arcm28\\zdot\\arcs8$ & 3.47 & 13.476 & --   &   0&.&6462   & E\\\\ 484 & BLG654.19 &  17980 & 17\\uph37\\upm19\\zdot\\ups57 & $-29\\arcd33\\arcm23\\zdot\\arcs8$ & 0.57 & 16.433 & --   &   7&.&062    & SP\\\\ 491 & BLG653.04 &  15499 & 17\\uph35\\upm43\\zdot\\ups46 & $-29\\arcd03\\arcm11\\zdot\\arcs8$ & 1.99 & 14.561 & --   &  16&.&10306  & P\\\\ 494 & BLG654.05 &  47692 & 17\\uph34\\upm58\\zdot\\ups69 & $-30\\arcd13\\arcm29\\zdot\\arcs0$ & 1.74 & 16.667 & --   &  --&&        & I\\\\ 495 & BLG654.22 &  76376 & 17\\uph34\\upm47\\zdot\\ups61 & $-29\\arcd35\\arcm10\\zdot\\arcs5$ & 2.09 & 16.623 & --   &   0&.&48629  & P\\\\ 497 & BLG654.14 &   1240 & 17\\uph34\\upm22\\zdot\\ups10 & $-30\\arcd05\\arcm05\\zdot\\arcs9$ & 0.20 & 14.057 & --   &   3&.&4965   & SP\\\\ 500 & BLG654.16 &  68256 & 17\\uph32\\upm39\\zdot\\ups40 & $-30\\arcd02\\arcm22\\zdot\\arcs7$ & 0.22 & 15.163 & --   &   1&.&8954   & P\\\\ 501 & BLG500.11 &  55525 & 17\\uph52\\upm52\\zdot\\ups63 & $-28\\arcd48\\arcm05\\zdot\\arcs4$ & 0.67 & 14.142 & 1.975 &  15&.&03759  & SP\\\\ 502 & BLG500.20 & 154130 & 17\\uph52\\upm23\\zdot\\ups57 & $-28\\arcd24\\arcm10\\zdot\\arcs9$ & 1.38 & 15.218 & 1.471 &   6&.&9252   & SP\\\\ 505 & BLG648.16 &  31001 & 17\\uph43\\upm53\\zdot\\ups28 & $-26\\arcd53\\arcm53\\zdot\\arcs7$ & 0.67 & 13.048 & --   &  11&.&147    & SP\\\\ 513 & BLG504.22 & 107805 & 17\\uph56\\upm46\\zdot\\ups04 & $-27\\arcd45\\arcm47\\zdot\\arcs6$ & 0.25 & 13.097 & 2.450 &  11&.&8343   & OSARG\\\\ 515 & BLG504.24 &     65 & 17\\uph55\\upm45\\zdot\\ups84 & $-27\\arcd58\\arcm13\\zdot\\arcs9$ & 3.14 & 14.572 & 2.383 &   4&.&95     & SP\\\\ 516 & BLG504.24 &  57283 & 17\\uph55\\upm34\\zdot\\ups69 & $-27\\arcd47\\arcm59\\zdot\\arcs7$ & 3.12 & 15.751 & --   &  33&.&333    & SRV\\\\ 522 & BLG500.17 & 158208 & 17\\uph54\\upm32\\zdot\\ups61 & $-28\\arcd29\\arcm19\\zdot\\arcs6$ & 1.96 & 12.971 & 5.007 & 147&.&3      & SRV\\\\ 526 & BLG500.01 & 144233 & 17\\uph54\\upm01\\zdot\\ups04 & $-29\\arcd00\\arcm46\\zdot\\arcs2$ & 3.71 & 15.429 & 2.513 & 135&.&1      & P\\\\ 532 & BLG500.21 &   2338 & 17\\uph52\\upm15\\zdot\\ups84 & $-28\\arcd34\\arcm15\\zdot\\arcs3$ & 0.73 & 12.890 & 1.873 &  44&.&25     & SP\\\\ 538 & BLG501.22 & 116664 & 17\\uph50\\upm54\\zdot\\ups10 & $-29\\arcd46\\arcm16\\zdot\\arcs0$ & 0.60 & 13.799 & 1.176 &   6&.&0533   & SP\\\\ 543 & BLG501.24 &  78370 & 17\\uph49\\upm32\\zdot\\ups70 & $-29\\arcd49\\arcm55\\zdot\\arcs9$ & 0.41 & 14.280 & 2.918 &  61&.&35     & SP\\\\ 556 & BLG648.29 &  41317 & 17\\uph46\\upm30\\zdot\\ups18 & $-26\\arcd22\\arcm04\\zdot\\arcs4$ & 0.20 & 14.675 & --   &  --&&        & I\\\\ 561 & BLG648.30 &   8035 & 17\\uph46\\upm07\\zdot\\ups68 & $-26\\arcd15\\arcm47\\zdot\\arcs1$ & 0.54 & 16.121 & --   &  12&.&4915   & E\\\\ 575 & BLG648.14 &  72850 & 17\\uph44\\upm59\\zdot\\ups28 & $-26\\arcd51\\arcm10\\zdot\\arcs2$ & 0.68 & 14.951 & --   &   6&.&5147   & SP\\\\ 594 & BLG675.18 &  76442 & 17\\uph42\\upm46\\zdot\\ups29 & $-27\\arcd13\\arcm53\\zdot\\arcs3$ & 0.16 & 16.472 & --   &  27&.&322    & SP\\\\ {\\bf 611} & BLG675.23 &  21477 & 17\\uph39\\upm50\\zdot\\ups50 & $-27\\arcd09\\arcm16\\zdot\\arcs5$ & 1.70 & 13.201 & --   &   1&.&3776   & P\\\\ {\\bf 611} & BLG675.23 &  21504 & 17\\uph39\\upm50\\zdot\\ups34 & $-27\\arcd09\\arcm17\\zdot\\arcs3$ & 2.98 & 14.793 & --   &  --&&        & I\\\\ 633 & BLG654.15 &  41782 & 17\\uph33\\upm18\\zdot\\ups72 & $-30\\arcd05\\arcm22\\zdot\\arcs3$ & 0.67 & 14.615 & --   &   0&.&73185  & P\\\\ 634 & BLG654.24 &  67281 & 17\\uph33\\upm16\\zdot\\ups65 & $-29\\arcd40\\arcm03\\zdot\\arcs2$ & 0.20 & 14.744 & --   &   0&.&76104  & SP\\\\ \\noalign{\\vskip3pt} \\hline} \\setcounter{table}{0} \\MakeTableSepp{r@{\\hspace{7pt}}c@{\\hspace{3pt}}r@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{5pt}}r@{\\hspace{0pt}}c@{\\hspace{0pt}}l@{\\hspace{5pt}}l} {12.5cm}{Continued} {\\hline \\noalign{\\vskip3pt} GBS               & OGLE-IV & OGLE-IV              &    RA    &   DEC    & $D$     & $I$   & $V-I$ & \\multicolumn{3}{c}{$P$}    & Remarks\\\\ \\multicolumn{1}{c}{ID} & Field   &\\multicolumn{1}{c}{No}& [2000.0] & [2000.0] & [\\arcs] & [mag] & [mag] & \\multicolumn{3}{c}{[days]} &\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} 637 & BLG654.15 &  75839 & 17\\uph33\\upm08\\zdot\\ups55 & $-29\\arcd57\\arcm27\\zdot\\arcs6$ & 1.21 & 13.429 & --   &   4&.&97018  & P\\\\ 642 & BLG504.14 & 192340 & 17\\uph55\\upm55\\zdot\\ups14 & $-28\\arcd01\\arcm58\\zdot\\arcs9$ & 1.12 & 14.924 & 1.948 &   2&.&924    & SP\\\\ 649 & BLG504.15 &  98233 & 17\\uph55\\upm24\\zdot\\ups65 & $-28\\arcd10\\arcm36\\zdot\\arcs8$ & 0.79 & 13.431 & 2.935 &  30&.&67485  & SP\\\\ 650 & BLG500.09 &  64937 & 17\\uph54\\upm08\\zdot\\ups51 & $-28\\arcd48\\arcm28\\zdot\\arcs8$ & 3.29 & 12.849 & 4.751 &  31&.&5      & SRV\\\\ 651 & BLG648.28 &  48495 & 17\\uph47\\upm03\\zdot\\ups70 & $-26\\arcd11\\arcm08\\zdot\\arcs8$ & 0.18 & 13.485 & --   &  --&&        & I\\\\ 657 & BLG504.29 &  15566 & 17\\uph57\\upm46\\zdot\\ups88 & $-27\\arcd33\\arcm09\\zdot\\arcs8$ & 0.39 & 16.280 & 1.562 &  16&.&58375  & P\\\\ 659 & BLG645.04 &  31747 & 17\\uph57\\upm38\\zdot\\ups37 & $-27\\arcd19\\arcm55\\zdot\\arcs2$ & 0.28 & 14.556 & --   &  74&.&627    & SP\\\\ 663 & BLG504.30 &  69448 & 17\\uph56\\upm30\\zdot\\ups96 & $-27\\arcd31\\arcm39\\zdot\\arcs9$ & 0.53 & 13.495 & 2.450 &  24&.&691    & SP\\\\ 666 & BLG504.23 & 118900 & 17\\uph56\\upm01\\zdot\\ups69 & $-27\\arcd44\\arcm48\\zdot\\arcs8$ & 0.49 & 14.104 & 1.522 &   2&.&6201   & E SP\\\\ 668 & BLG504.14 & 188107 & 17\\uph55\\upm55\\zdot\\ups86 & $-28\\arcd05\\arcm55\\zdot\\arcs5$ & 0.41 & 14.147 & 2.669 & 106&.&4      & SP\\\\ 674 & BLG504.24 &  74705 & 17\\uph55\\upm19\\zdot\\ups78 & $-27\\arcd56\\arcm32\\zdot\\arcs5$ & 1.07 & 14.868 & --   &  24&.&73     & OSARG\\\\ 676 & BLG500.17 & 159235 & 17\\uph54\\upm32\\zdot\\ups45 & $-28\\arcd28\\arcm13\\zdot\\arcs4$ & 3.73 & 17.241 & 2.748 &   0&.&30728  & RRLYR\\\\ 679 & BLG500.18 &  50697 & 17\\uph54\\upm13\\zdot\\ups98 & $-28\\arcd33\\arcm41\\zdot\\arcs6$ & 3.68 & 15.111 & 3.425 &  --&&        & I\\\\ 681 & BLG500.09 & 128064 & 17\\uph54\\upm05\\zdot\\ups63 & $-28\\arcd47\\arcm32\\zdot\\arcs6$ & 0.04 & 13.865 & 0.826 &   2&.&0736   & SP\\\\ 684 & BLG500.09 & 136274 & 17\\uph54\\upm00\\zdot\\ups38 & $-28\\arcd43\\arcm32\\zdot\\arcs4$ & 2.47 & 17.369 & 2.566 &   2&.&5864   & E SP\\\\ 685 & BLG500.09 & 128146 & 17\\uph53\\upm56\\zdot\\ups51 & $-28\\arcd46\\arcm58\\zdot\\arcs3$ & 0.58 & 14.760 & 1.345 &   2&.&7894   & SP\\\\ 688 & BLG500.19 &    348 & 17\\uph53\\upm42\\zdot\\ups30 & $-28\\arcd35\\arcm31\\zdot\\arcs8$ & 0.63 & 16.640 & 3.175 &   8&.&231    & SP\\\\ 689 & BLG500.19 &  17579 & 17\\uph53\\upm39\\zdot\\ups84 & $-28\\arcd29\\arcm31\\zdot\\arcs6$ & 0.36 & 13.825 & 1.525 & 189& &       & P\\\\ 691 & BLG500.19 & 125229 & 17\\uph53\\upm24\\zdot\\ups21 & $-28\\arcd27\\arcm10\\zdot\\arcs8$ & 1.20 & 19.629 & --   &   2&.&794    & P\\\\ 692 & BLG500.02 & 163813 & 17\\uph53\\upm05\\zdot\\ups18 & $-29\\arcd13\\arcm00\\zdot\\arcs6$ & 1.33 & 14.609 & 3.796 &  93&.&3      & P\\\\ 695 & BLG501.28 &  64217 & 17\\uph52\\upm40\\zdot\\ups19 & $-29\\arcd20\\arcm57\\zdot\\arcs8$ & 0.26 & 15.703 & 2.757 &   0&.&30988  & P\\\\ 701 & BLG501.29 & 138642 & 17\\uph51\\upm41\\zdot\\ups10 & $-29\\arcd18\\arcm55\\zdot\\arcs2$ & 0.77 & 13.531 & 1.838 &  11&.&9936   & E SP\\\\ 702 & BLG500.12 & 122438 & 17\\uph51\\upm39\\zdot\\ups41 & $-28\\arcd52\\arcm43\\zdot\\arcs5$ & 0.44 & 15.540 & 1.560 &   1&.&50011  & E\\\\ {\\bf 705} & BLG501.30 &  35398 & 17\\uph51\\upm16\\zdot\\ups15 & $-29\\arcd23\\arcm44\\zdot\\arcs4$ & 0.87 & 15.975 & 1.579 &  --&&        & I \\\\ {\\bf 705} & BLG501.30 &  35400 & 17\\uph51\\upm15\\zdot\\ups91 & $-29\\arcd23\\arcm42\\zdot\\arcs8$ & 2.76 & 15.400 & --   & 326&.&8      & MIRA\\\\ 711 & BLG501.31 &  34198 & 17\\uph50\\upm32\\zdot\\ups49 & $-29\\arcd21\\arcm00\\zdot\\arcs1$ & 0.68 & 15.649 & 2.322 &  13&.&73     & SP\\\\ 712 & BLG501.14 & 108924 & 17\\uph50\\upm14\\zdot\\ups77 & $-30\\arcd03\\arcm19\\zdot\\arcs5$ & 0.76 & 15.556 & 2.485 &   6&.&84     & SP\\\\ 714 & BLG501.15 &  11770 & 17\\uph49\\upm59\\zdot\\ups31 & $-29\\arcd57\\arcm11\\zdot\\arcs1$ & 0.71 & 18.874 & --   &  --&&        & CV\\\\ 715 & BLG501.24 &  18825 & 17\\uph49\\upm59\\zdot\\ups15 & $-29\\arcd34\\arcm38\\zdot\\arcs1$ & 0.79 & 14.836 & --   &  42&.&735    & OSARG\\\\ 717 & BLG501.24 &  35909 & 17\\uph49\\upm50\\zdot\\ups32 & $-29\\arcd42\\arcm30\\zdot\\arcs4$ & 2.92 & 18.600 & --   &   0&.&54767  & RRLYR$^*$\\\\ 718 & BLG501.07 &  45474 & 17\\uph49\\upm48\\zdot\\ups54 & $-30\\arcd12\\arcm21\\zdot\\arcs0$ & 0.65 & 15.335 & 1.437 &   5&.&824    & P\\\\ 721 & BLG534.32 & 107583 & 17\\uph49\\upm20\\zdot\\ups74 & $-30\\arcd28\\arcm05\\zdot\\arcs9$ & 0.34 & 15.633 & --   &   2&.&66207  & E\\\\ {\\it   724}& {\\it BUL\\_SC44} & {\\it 12986} & {\\it 17\\uph48\\upm54\\zdot\\ups26} & ${\\it -30\\arcd18\\arcm39\\zdot\\arcs5}$ & {\\it 0.51} & {\\it 12.827} & {\\it 1.104} &   {\\it 0}&.&{\\it 32200}  & {\\it E}\\\\ 742 & BLG648.29 &  69672 & 17\\uph46\\upm19\\zdot\\ups94 & $-26\\arcd12\\arcm28\\zdot\\arcs7$ & 0.11 & 14.430 & --   &   0&.&25321  & P\\\\ 774 & BLG675.17 &   4519 & 17\\uph44\\upm01\\zdot\\ups34 & $-27\\arcd16\\arcm46\\zdot\\arcs8$ & 3.57 & 14.530 & --   &  43&.&47826  & P\\\\ 778 & BLG648.16 &  78697 & 17\\uph43\\upm35\\zdot\\ups43 & $-26\\arcd48\\arcm57\\zdot\\arcs1$ & 0.44 & 17.822 & --   &   6&.&079    & P\\\\ 794 & BLG675.03 &  22881 & 17\\uph41\\upm42\\zdot\\ups38 & $-27\\arcd58\\arcm30\\zdot\\arcs3$ & 0.78 & 19.209 & --   &   0&.&11786  & E?\\\\ 796 & BLG675.03 &  73380 & 17\\uph41\\upm27\\zdot\\ups62 & $-27\\arcd58\\arcm39\\zdot\\arcs2$ & 0.48 & 14.229 & --   &   1&.&06571  & E\\\\ 800 & BLG675.04 &  85070 & 17\\uph40\\upm51\\zdot\\ups59 & $-27\\arcd45\\arcm59\\zdot\\arcs7$ & 0.32 & 13.954 & --   & 213& &       & MIRA\\\\ 809 & BLG675.22 &  60737 & 17\\uph40\\upm17\\zdot\\ups79 & $-27\\arcd06\\arcm07\\zdot\\arcs8$ & 1.26 & 14.581 & --   &  35&.&7      & P\\\\ 811 & BLG675.14 &  28661 & 17\\uph39\\upm48\\zdot\\ups50 & $-27\\arcd24\\arcm16\\zdot\\arcs0$ & 0.59 & 14.668 & --   &  43&.&3      & P\\\\ 819 & BLG675.07 &  56138 & 17\\uph38\\upm56\\zdot\\ups46 & $-27\\arcd48\\arcm33\\zdot\\arcs4$ & 0.36 & 16.978 & --   &  35& &       & SP\\\\ 829 & BLG654.26 &  41915 & 17\\uph37\\upm38\\zdot\\ups06 & $-29\\arcd21\\arcm44\\zdot\\arcs8$ & 0.42 & 15.767 & --   &  --&&        & I\\\\ 830 & BLG653.26 &  52494 & 17\\uph37\\upm29\\zdot\\ups94 & $-28\\arcd07\\arcm59\\zdot\\arcs5$ & 0.58 & 13.127 & --   &   8&.&288    & E\\\\ 837 & BLG653.19 &  61176 & 17\\uph36\\upm47\\zdot\\ups64 & $-28\\arcd22\\arcm21\\zdot\\arcs7$ & 2.41 & 15.714 & --   &  15&.&18     & P\\\\ 843 & BLG654.03 &  59545 & 17\\uph36\\upm05\\zdot\\ups70 & $-30\\arcd18\\arcm16\\zdot\\arcs8$ & 3.27 & 14.403 & --   &  49&.&5      & OSARG\\\\ 845 & BLG654.29 &   2771 & 17\\uph35\\upm47\\zdot\\ups48 & $-29\\arcd29\\arcm15\\zdot\\arcs0$ & 0.08 & 16.443 & --   &   3&.&158    & P\\\\ 847 & BLG654.12 &  30313 & 17\\uph35\\upm37\\zdot\\ups78 & $-29\\arcd56\\arcm01\\zdot\\arcs3$ & 1.00 & 14.470 & --   &  17&.&212    & P\\\\ 848 & BLG654.04 &  55881 & 17\\uph35\\upm33\\zdot\\ups43 & $-30\\arcd17\\arcm44\\zdot\\arcs0$ & 0.42 & 18.294 & --   &   7&.&4405   & SP\\\\ 852 & BLG654.13 &  32681 & 17\\uph35\\upm02\\zdot\\ups21 & $-30\\arcd03\\arcm54\\zdot\\arcs7$ & 0.93 & 14.742 & --   &   0&.&29857  & P\\\\ 853 & BLG654.13 &  74975 & 17\\uph34\\upm46\\zdot\\ups94 & $-29\\arcd56\\arcm31\\zdot\\arcs6$ & 0.64 & 16.118 & --   &   0&.&91405  & E\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} \\multicolumn{12}{l}{$^*$ OGLE-BLG-RRLYR-03686 (Soszy\u00f1ski \\etal 2011a)} } \\setcounter{table}{0} \\MakeTableSepp{r@{\\hspace{7pt}}c@{\\hspace{3pt}}r@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{5pt}}r@{\\hspace{0pt}}c@{\\hspace{0pt}}l@{\\hspace{5pt}}l} {12.5cm}{Concluded} {\\hline \\noalign{\\vskip3pt} GBS               & OGLE-IV & OGLE-IV              &    RA    &   DEC    & $D$     & $I$   & $V-I$ & \\multicolumn{3}{c}{$P$}    & Remarks\\\\ \\multicolumn{1}{c}{ID} & Field   &\\multicolumn{1}{c}{No}& [2000.0] & [2000.0] & [\\arcs] & [mag] & [mag] & \\multicolumn{3}{c}{[days]} &\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} 858 & BLG654.24 &  63566 & 17\\uph33\\upm26\\zdot\\ups90 & $-29\\arcd42\\arcm11\\zdot\\arcs7$ & 0.27 & 15.403 & --   &  11&.&174    & SP\\\\ 869 & BLG648.13 &  45934 & 17\\uph45\\upm36\\zdot\\ups79 & $-26\\arcd51\\arcm32\\zdot\\arcs7$ & 0.31 & 14.130 & --   &  23&.&81     & SP\\\\ 877 & BLG504.24 & 106420 & 17\\uph55\\upm14\\zdot\\ups65 & $-27\\arcd59\\arcm01\\zdot\\arcs4$ & 0.36 & 12.763 & 2.472 &  27&.&933    & SP\\\\ 878 & BLG501.24 &   8191 & 17\\uph50\\upm01\\zdot\\ups72 & $-29\\arcd42\\arcm56\\zdot\\arcs7$ & 0.71 & 15.848 & 1.646 &   0&.&64616  & E\\\\ 895 & BLG654.23 &  73640 & 17\\uph34\\upm05\\zdot\\ups41 & $-29\\arcd42\\arcm43\\zdot\\arcs8$ & 0.81 & 15.979 & --   &   0&.&42973  & E\\\\ 896 & BLG645.03 &  63999 & 17\\uph58\\upm12\\zdot\\ups28 & $-27\\arcd21\\arcm39\\zdot\\arcs3$ & 0.77 & 13.707 & --   &   3&.&8534   & E\\\\ 898 & BLG504.29 &  96872 & 17\\uph57\\upm28\\zdot\\ups81 & $-27\\arcd33\\arcm10\\zdot\\arcs7$ & 0.60 & 14.798 & 3.099 &  44&.&84     & SP\\\\ 899 & BLG504.29 & 118157 & 17\\uph57\\upm18\\zdot\\ups94 & $-27\\arcd36\\arcm08\\zdot\\arcs7$ & 0.34 & 12.976 & 3.020 &  49&.&751    & SP\\\\ 901 & BLG504.30 &  59047 & 17\\uph56\\upm51\\zdot\\ups19 & $-27\\arcd23\\arcm50\\zdot\\arcs6$ & 0.75 & 13.094 & 0.829 &   0&.&39967  & E\\\\ 906 & BLG504.31 &  26678 & 17\\uph56\\upm16\\zdot\\ups81 & $-27\\arcd26\\arcm34\\zdot\\arcs2$ & 0.55 & 13.798 & 2.126 &   0&.&29347  & SP\\\\ 907 & BLG504.23 &  59042 & 17\\uph56\\upm11\\zdot\\ups47 & $-27\\arcd52\\arcm33\\zdot\\arcs5$ & 3.43 & 15.710 & --   & 172&.&4      & SRV\\\\ 908 & BLG504.31 &  22399 & 17\\uph56\\upm10\\zdot\\ups29 & $-27\\arcd32\\arcm58\\zdot\\arcs1$ & 1.08 & 17.148 & --   &  67&.&1      & SRV\\\\ 909 & BLG504.31 &  20360 & 17\\uph56\\upm09\\zdot\\ups94 & $-27\\arcd35\\arcm17\\zdot\\arcs5$ & 0.42 & 15.969 & 2.516 &   4&.&0967   & SP\\\\ 910 & BLG504.23 &  90667 & 17\\uph56\\upm07\\zdot\\ups52 & $-27\\arcd57\\arcm26\\zdot\\arcs4$ & 0.23 & 15.288 & 1.947 &   1&.&24425  & SP\\\\ 912 & BLG504.31 &  42181 & 17\\uph56\\upm03\\zdot\\ups86 & $-27\\arcd28\\arcm33\\zdot\\arcs6$ & 0.62 & 15.458 & 2.346 &   0&.&37699  & P\\\\ 914 & BLG504.31 &  33921 & 17\\uph55\\upm59\\zdot\\ups55 & $-27\\arcd37\\arcm44\\zdot\\arcs9$ & 0.65 & 13.769 & 0.833 &  --&&        & I\\\\ 915 & BLG504.31 &  58984 & 17\\uph55\\upm58\\zdot\\ups61 & $-27\\arcd29\\arcm56\\zdot\\arcs8$ & 2.29 & 16.630 & 1.647 &   1&.&1609   & P\\\\ 924 & BLG500.17 &  36873 & 17\\uph55\\upm00\\zdot\\ups48 & $-28\\arcd19\\arcm56\\zdot\\arcs5$ & 0.20 & 13.777 & 1.991 &  77&.&5      & P\\\\ 926 & BLG504.16 &  57832 & 17\\uph54\\upm49\\zdot\\ups47 & $-28\\arcd07\\arcm36\\zdot\\arcs6$ & 0.42 & 17.195 & 2.210 &   3&.&787    & P\\\\ 929 & BLG500.17 & 144211 & 17\\uph54\\upm30\\zdot\\ups83 & $-28\\arcd33\\arcm35\\zdot\\arcs5$ & 0.19 & 16.227 & 1.695 &   0&.&59379  & SP\\\\ 930 & BLG500.17 & 151398 & 17\\uph54\\upm29\\zdot\\ups31 & $-28\\arcd31\\arcm52\\zdot\\arcs8$ & 2.99 & 16.716 & 1.665 &  --&&        & I\\\\ 933 & BLG500.09 & 120810 & 17\\uph54\\upm05\\zdot\\ups77 & $-28\\arcd48\\arcm23\\zdot\\arcs4$ & 2.27 & 15.607 & 3.306 &  --&&        & I\\\\ 938 & BLG500.09 & 190630 & 17\\uph53\\upm52\\zdot\\ups05 & $-28\\arcd44\\arcm52\\zdot\\arcs7$ & 1.33 & 13.735 & 0.786 &   2&.&42096  & E\\\\ 940 & BLG500.19 &  52204 & 17\\uph53\\upm30\\zdot\\ups09 & $-28\\arcd32\\arcm47\\zdot\\arcs7$ & 0.39 & 15.330 & 1.337 &  --&&        & I\\\\ 944 & BLG500.28 &     13 & 17\\uph52\\upm54\\zdot\\ups34 & $-28\\arcd16\\arcm29\\zdot\\arcs9$ & 2.48 & 14.334 & 5.269 &  42&.&02     & SRV\\\\ 946 & BLG500.11 &  55525 & 17\\uph52\\upm52\\zdot\\ups63 & $-28\\arcd48\\arcm05\\zdot\\arcs4$ & 3.22 & 14.142 & 1.975 &  15&.&03759  & SP\\\\ 947 & BLG500.03 &  89767 & 17\\uph52\\upm45\\zdot\\ups16 & $-28\\arcd59\\arcm15\\zdot\\arcs1$ & 3.42 & 15.918 & 5.493 &  14&.&881    & OSARG\\\\ 948 & BLG500.03 & 100501 & 17\\uph52\\upm41\\zdot\\ups86 & $-29\\arcd12\\arcm52\\zdot\\arcs6$ & 0.36 & 14.925 & 1.430 &   1&.&9577   & SP\\\\ 953 & BLG500.21 &  34612 & 17\\uph52\\upm03\\zdot\\ups78 & $-28\\arcd33\\arcm51\\zdot\\arcs1$ & 0.34 & 15.901 & 1.534 &   1&.&0989   & SP\\\\ 955 & BLG500.12 & 106387 & 17\\uph51\\upm51\\zdot\\ups42 & $-28\\arcd42\\arcm46\\zdot\\arcs5$ & 0.35 & 16.386 & --   &  --&&        & I\\\\ {\\it   964}& {\\it BUL\\_SC5} & {\\it 172053} & {\\it 17\\uph50\\upm15\\zdot\\ups31} & ${\\it -29\\arcd39\\arcm13\\zdot\\arcs8}$ & {\\it 0.19} & {\\it 12.289} & {\\it 1.681} &  {\\it 81}&.&{\\it 97} & {\\it P}\\\\ 966 & BLG501.15 &  49578 & 17\\uph49\\upm42\\zdot\\ups67 & $-30\\arcd02\\arcm44\\zdot\\arcs5$ & 1.96 & 16.881 & 3.170 &   9&.&662    & SP\\\\ 969 & BLG501.15 &  71683 & 17\\uph49\\upm32\\zdot\\ups21 & $-30\\arcd03\\arcm50\\zdot\\arcs4$ & 0.60 & 18.335 & --   &  32&.&154    & P\\\\ 970 & BLG501.24 &  80876 & 17\\uph49\\upm28\\zdot\\ups07 & $-29\\arcd46\\arcm29\\zdot\\arcs1$ & 0.70 & 12.504 & 0.970 &  --&&        & I\\\\ 972 & BLG534.32 & 100700 & 17\\uph49\\upm23\\zdot\\ups67 & $-30\\arcd32\\arcm16\\zdot\\arcs2$ & 0.30 & 12.755 & 0.921 &   0&.&334094 & E\\\\ 973 & BLG501.16 &     69 & 17\\uph49\\upm21\\zdot\\ups61 & $-30\\arcd07\\arcm39\\zdot\\arcs2$ & 0.45 & 16.672 & 2.838 &  10&.&352    & SP\\\\ 978 & BLG501.25 &  42282 & 17\\uph48\\upm59\\zdot\\ups54 & $-29\\arcd39\\arcm40\\zdot\\arcs0$ & 2.32 & 16.675 & 1.708 &   1&.&05765  & P\\\\ 981 & BLG501.16 &  49689 & 17\\uph48\\upm54\\zdot\\ups94 & $-29\\arcd54\\arcm16\\zdot\\arcs7$ & 0.51 & 15.823 & 1.335 &  91&.&7      & P\\\\ 1019 & BLG648.13 &  46396 & 17\\uph45\\upm48\\zdot\\ups53 & $-26\\arcd50\\arcm50\\zdot\\arcs5$ & 0.18 & 16.406 & --   &   0&.&35143  & E\\\\ 1033 & BLG648.31 &  60638 & 17\\uph45\\upm02\\zdot\\ups19 & $-26\\arcd14\\arcm49\\zdot\\arcs4$ & 0.23 & 13.104 & --   &  54&.&07799  & SP\\\\ 1070 & BLG675.26 &  56707 & 17\\uph42\\upm58\\zdot\\ups73 & $-26\\arcd55\\arcm01\\zdot\\arcs4$ & 0.39 & 14.560 & --   &  47&.&17     & SP\\\\ 1071 & BLG675.01 &  49220 & 17\\uph42\\upm57\\zdot\\ups12 & $-27\\arcd46\\arcm30\\zdot\\arcs2$ & 0.55 & 15.767 & --   &  37&.&175    & P\\\\ 1074 & BLG675.26 &  74928 & 17\\uph42\\upm52\\zdot\\ups61 & $-26\\arcd57\\arcm14\\zdot\\arcs1$ & 0.96 & 15.156 & --   &   3&.&1616   & E\\\\ 1078 & BLG675.19 &  37341 & 17\\uph42\\upm24\\zdot\\ups03 & $-27\\arcd11\\arcm57\\zdot\\arcs8$ & 2.61 & 16.536 & --   &  --&&        & I\\\\ 1086 & BLG675.03 &  34167 & 17\\uph41\\upm46\\zdot\\ups47 & $-27\\arcd49\\arcm54\\zdot\\arcs7$ & 0.03 & 16.526 & --   &  11&.&768    & E\\\\ 1090 & BLG675.11 & 107190 & 17\\uph41\\upm25\\zdot\\ups07 & $-27\\arcd31\\arcm43\\zdot\\arcs6$ & 0.67 & 14.827 & --   &  49&.&26     & SP\\\\ 1104 & BLG675.21 & 104241 & 17\\uph40\\upm40\\zdot\\ups67 & $-27\\arcd08\\arcm07\\zdot\\arcs9$ & 3.88 & 15.119 & --   & 109&.&89     & P\\\\ 1110 & BLG675.13 &  40278 & 17\\uph40\\upm17\\zdot\\ups69 & $-27\\arcd30\\arcm57\\zdot\\arcs1$ & 0.44 & 13.617 & --   &  30&.&612    & E\\\\ 1112 & BLG675.22 &  79334 & 17\\uph40\\upm12\\zdot\\ups92 & $-27\\arcd12\\arcm35\\zdot\\arcs7$ & 1.26 & 14.275 & --   &  17&.&699    & SP\\\\ 1151 & BLG653.01 &   1611 & 17\\uph37\\upm55\\zdot\\ups27 & $-29\\arcd10\\arcm10\\zdot\\arcs1$ & 2.67 & 14.616 & --   &  78&.&125    & SRV\\\\ {\\bf 1155} & BLG654.26 &  59596 & 17\\uph37\\upm31\\zdot\\ups10 & $-29\\arcd24\\arcm28\\zdot\\arcs8$ & 0.38 & 13.476 & --   &   0&.&6462   & E\\\\ {\\bf 1155} & BLG654.26 &  59618 & 17\\uph37\\upm31\\zdot\\ups33 & $-29\\arcd24\\arcm28\\zdot\\arcs9$ & 3.35 & 15.531 & --   &  56&.&8      & SRV\\\\ \\noalign{\\vskip3pt} \\hline} \\setcounter{table}{0} \\MakeTable{r@{\\hspace{7pt}}c@{\\hspace{3pt}}r@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{7pt}}c@{\\hspace{5pt}}r@{\\hspace{0pt}}c@{\\hspace{0pt}}l@{\\hspace{5pt}}l} {12.5cm}{Concluded} {\\hline \\noalign{\\vskip3pt} GBS               & OGLE-IV & OGLE-IV              &    RA    &   DEC    & $D$     & $I$   & $V-I$ & \\multicolumn{3}{c}{$P$}    & Remarks\\\\ \\multicolumn{1}{c}{ID} & Field   &\\multicolumn{1}{c}{No}& [2000.0] & [2000.0] & [\\arcs] & [mag] & [mag] & \\multicolumn{3}{c}{[days]} &\\\\ \\noalign{\\vskip3pt} \\hline 1161 & BLG653.19 &  67406 & 17\\uph36\\upm45\\zdot\\ups56 & $-28\\arcd34\\arcm14\\zdot\\arcs8$ & 3.15 & 15.511 & --   &  94&.&3      & SRV\\\\ 1165 & BLG653.20 &  10600 & 17\\uph36\\upm33\\zdot\\ups67 & $-28\\arcd30\\arcm03\\zdot\\arcs9$ & 0.14 & 13.703 & --   &  74&.&1      & SRV\\\\ 1168 & BLG654.03 &  26666 & 17\\uph36\\upm27\\zdot\\ups84 & $-30\\arcd12\\arcm14\\zdot\\arcs1$ & 0.52 & 13.823 & --   &  65&.&79     & P\\\\ 1179 & BLG654.11 &  52361 & 17\\uph36\\upm06\\zdot\\ups21 & $-29\\arcd55\\arcm40\\zdot\\arcs5$ & 2.14 & 15.369 & --   &  --&&        & I\\\\ 1185 & BLG653.12 &   9604 & 17\\uph35\\upm48\\zdot\\ups65 & $-28\\arcd47\\arcm09\\zdot\\arcs7$ & 0.94 & 14.880 & --   &   2&.&1497   & E\\\\ 1187 & BLG653.04 &   7005 & 17\\uph35\\upm45\\zdot\\ups73 & $-29\\arcd08\\arcm02\\zdot\\arcs7$ & 2.71 & 15.012 & --   &  16&.&92     & OSARG\\\\ 1194 & BLG654.04 &  62415 & 17\\uph35\\upm27\\zdot\\ups50 & $-30\\arcd14\\arcm01\\zdot\\arcs3$ & 2.51 & 16.015 & --   &   1&.&91022  & SP\\\\ 1208 & BLG654.06 &  16127 & 17\\uph34\\upm23\\zdot\\ups71 & $-30\\arcd17\\arcm05\\zdot\\arcs1$ & 0.90 & 14.935 & --   &  24&.&876    & SP\\\\ 1228 & BLG500.03 & 188094 & 17\\uph52\\upm31\\zdot\\ups10 & $-28\\arcd58\\arcm03\\zdot\\arcs9$ & 1.58 & 14.929 & 2.458 &   3&.&4235   & P\\\\ 1229 & BLG500.12 & 106361 & 17\\uph51\\upm56\\zdot\\ups83 & $-28\\arcd41\\arcm45\\zdot\\arcs7$ & 0.34 & 14.649 & 1.770 &  --&&        & I\\\\ \\noalign{\\vskip3pt} \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} \\multicolumn{12}{p{12.5cm}}{ $\\bullet$ Italic font entry highlights the counterpart candidates selected from the OGLE-II databases.\\newline $\\bullet$ Boldface GBS number highlights the cases when two variable stars are located in the search region.\\newline $\\bullet$ Variability type abbreviations: SP~--~spotted star, OSARG/SRV/MIRA -- pulsating giant, E~--~eclipsing system, P -- periodic system, RRLYR~--~RR~Lyr pulsating star, CV~--~cataclysmic variable, I~--~irregular star, AGN -- active galactic nucleus.\\newline $\\bullet$ The same optical counterpart candidates for two close GBS X-ray sources: GBS \\#16/GBS \\#515, GBS \\#303/GBS \\#304, GBS \\#483/GBS \\#1155, GBS \\#501/GBS \\#946. } } It should be noted that in the case of five X-ray sources from the GBS catalog more than one variable star were found within our search radius. The best example here is the X-ray source GBS \\#19 (Fig.~7). A variable red giant and fainter eclipsing system are detected very close to the X-ray position. Only further, preferably spectroscopic, observations may solve the cross-identification ambiguity in this case.  However, it may turn out that both these stars are X-ray sources. A one day long flare of $\\approx 1$~mag amplitude observed on ${\\rm HJD=2\\,455\\,349.7}$ and ${\\rm HJD=2\\,456\\,071.7}$ in the eclipsing star may indicate stellar activity. On the other hand the light curve of the red giant resembles that of other X-ray binary systems.  There are also four cases where two separate X-rays sources from the GBS catalog point very close to each other (\\eg\\ \\#16 and \\#515 are separated by only 3\\zdot\\arcs3). In such cases the same optical variable stars are selected as potential counterparts to these X-ray sources. Additional observations can resolve the issue of the real connections of the sources in such cases.  \\begin{figure}[htb] \\centerline{\\includegraphics[width=8.2cm, bb=110 50 510 310]{fig8.ps}} \\FigCap{Histogram of the angular distances in the sky between the OGLE and X-ray position of the optical counterpart candidates. The bin size is 0\\zdot\\arcs1.} \\end{figure} Fig.~8 shows the histogram of the angular distances between the OGLE equatorial coordinates of the counterpart candidates and their X-ray positions in 0\\zdot\\arcs1 bins. It is clearly seen that for the vast majority of objects the coincidence of positions is better than 1\\arcs. In the remaining cases it is possible that the X-ray coordinates are less accurate or a variable optical candidate is aligned by chance on the sky with the corresponding X-ray source. Conservative selection of the 15 pixel (3\\zdot\\arcs9) radius ensures high completeness of our search.  \\MakeTable{rcrccccc}{12.5cm}{Possible bright non-variable optical counterparts to the X-ray sources detected by the Galactic Bulge Survey} {\\hline \\noalign{\\vskip3pt} GBS &  OGLE-IV  &OGLE-IV &  RA      & DEC      &  $D$    & $I$   & $V-I$ \\\\ ID  &  Field    &\\multicolumn{1}{c}{No} & [2000.0] & [2000.0] & [\\arcs] & [mag] & [mag]\\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} 88 & BLG654.24 &  96086 & 17\\uph33\\upm08\\zdot\\ups04 & $-29\\arcd37\\arcm52\\zdot\\arcs7$ & 0.36 & 15.423 & --\\\\ 149 & BLG675.27 &  31995 & 17\\uph42\\upm26\\zdot\\ups78 & $-26\\arcd47\\arcm31\\zdot\\arcs3$ & 0.57 & 12.636 & --\\\\ 170 & BLG654.18 &  34390 & 17\\uph37\\upm55\\zdot\\ups18 & $-29\\arcd35\\arcm40\\zdot\\arcs3$ & 0.13 & 12.980 & --\\\\ 294 & BLG504.25 &  67635 & 17\\uph54\\upm35\\zdot\\ups37 & $-27\\arcd44\\arcm59\\zdot\\arcs4$ & 0.49 & 14.041 & 1.783\\\\ 393 & BLG654.29 &  66612 & 17\\uph35\\upm30\\zdot\\ups83 & $-29\\arcd18\\arcm31\\zdot\\arcs9$ & 0.23 & 14.145 & --\\\\ 403 & BLG500.02 & 184262 & 17\\uph53\\upm06\\zdot\\ups53 & $-29\\arcd05\\arcm30\\zdot\\arcs1$ & 0.33 & 14.384 & 1.055\\\\ 454 & BLG648.23 &   8019 & 17\\uph45\\upm25\\zdot\\ups15 & $-26\\arcd39\\arcm08\\zdot\\arcs0$ & 0.49 & 12.490 & --\\\\ 631 & BLG653.05 &  53272 & 17\\uph34\\upm43\\zdot\\ups92 & $-29\\arcd09\\arcm49\\zdot\\arcs7$ & 0.60 & 13.521 & --\\\\ 646 & BLG654.30 &  12088 & 17\\uph35\\upm06\\zdot\\ups17 & $-29\\arcd21\\arcm12\\zdot\\arcs1$ & 0.36 & 13.802 & --\\\\ 662 & BLG504.22 & 107804 & 17\\uph56\\upm42\\zdot\\ups18 & $-27\\arcd45\\arcm54\\zdot\\arcs3$ & 0.10 & 13.144 & 1.859\\\\ 675 & BLG500.17 & 150996 & 17\\uph54\\upm34\\zdot\\ups00 & $-28\\arcd32\\arcm19\\zdot\\arcs7$ & 0.36 & 13.360 & 0.906\\\\ 773 & BLG648.16 &  12494 & 17\\uph44\\upm01\\zdot\\ups99 & $-26\\arcd50\\arcm49\\zdot\\arcs7$ & 0.75 & 13.285 & --\\\\ 880 & BLG501.16 &  47309 & 17\\uph49\\upm00\\zdot\\ups89 & $-29\\arcd54\\arcm32\\zdot\\arcs8$ & 0.22 & 13.102 & 0.794\\\\ 967 & BLG534.32 &  76738 & 17\\uph49\\upm37\\zdot\\ups15 & $-30\\arcd32\\arcm57\\zdot\\arcs8$ & 0.34 & 13.074 & 0.719\\\\ 1045 & BLG648.32 &  45146 & 17\\uph44\\upm29\\zdot\\ups95 & $-26\\arcd19\\arcm49\\zdot\\arcs5$ & 0.29 & 15.861 & --\\\\ 1049 & BLG648.16 &   1014 & 17\\uph44\\upm08\\zdot\\ups94 & $-26\\arcd59\\arcm58\\zdot\\arcs6$ & 0.35 & 15.554 & --\\\\ 1123 & BLG675.06 &  51526 & 17\\uph39\\upm38\\zdot\\ups91 & $-27\\arcd43\\arcm56\\zdot\\arcs5$ & 0.15 & 16.286 & --\\\\ 1200 & BLG654.22 &  10352 & 17\\uph35\\upm13\\zdot\\ups17 & $-29\\arcd40\\arcm23\\zdot\\arcs8$ & 0.27 & 14.655 & --\\\\ 1209 & BLG654.31 &  13876 & 17\\uph34\\upm23\\zdot\\ups62 & $-29\\arcd21\\arcm55\\zdot\\arcs0$ & 0.67 & 13.471 & --\\\\ \\noalign{\\vskip3pt} \\hline} Finally, it is worth noticing that in the case of 19 X-ray sources where we did not find any optically variable objects in their neighborhood, the position of the X-ray source points with a high accuracy -- better than 0\\zdot\\arcs75 -- to a well separated star brighter than 16.5~mag. Therefore, it is very likely that these stars may also be potential optical counterparts. As can be seen from the light curves in the OGLE X-ray optical counterparts monitoring system, XROM (Udalski 2008), many optical counterparts to the X-ray sources can be optically quiet for years. Table~2 lists these potential non-variable optical counterparts to the X-ray sources.  Table~1 can be a good starting point for further follow-up observations of the selected optical counterparts. In particular extracting new low-mass and massive X-ray binaries would be of great importance for further studies of these objects in the Galactic center environment. We plan to include a subsample of the most interesting objects presented here to the list of objects monitored regularly by the OGLE XROM system (Udalski 2008). This can facilitate follow-up observations providing information on the current optical behavior of the observed objects."
    },
    "1208/1208.1435_arXiv.txt": {
        "abstract": "We analyse a feature that is detectable at the differential phases of high-resolution spectro-interferometry of Be stars. The origin of this feature (dubbed CQE-PS, Central Quasi Emission Phase Signature) lies in the differential absorption of photospheric radiation by the circumstellar disk, which is spectroscopically detected as a CQE line profile in shell stars. This phenomenon has great diagnostic potential for Be star disks, revealing properties of these disks on the scale of a few stellar radii. ",
        "introduction": "\\label{intro}  Classical Be stars are well known to be characterized by having gaseous circumstellar disks that are fed from mass lost from their rapidly rotating central stars \\citep{por03}. In the past two decades, optical and infrared interferometry played a key role in the study of these astrophysical systems. For instance, the pioneering measurements of \\citet{qui97} provided strong evidence that the circumstellar matter around Be stars is distributed in a geometrically thin structure (i.e., a disk). More recently, detailed measurements allowed to demonstrate that Be disks rotate in a Keplerian fashion \\citep{mei07,kra11,whe12}.  The strong developments in the observations were accompanied by significant progress in our theoretical understanding of these systems. Recently, a consensus is emerging that the viscous decretion disk model (VDDM), first suggested by \\cite{lee91} and further developed by Bjorkman (1997), and Okazaki (2001), among others, is the best candidate to explain the observed properties of Be disks \\citep{car11,mcg11}. This model uses the angular momentum transport by turbulent viscosity to lift material that was ejected by the star into higher orbits, thereby causing the disk to grow in size. A theoretical result of this model is that material is in Keplerian rotation throughout the disk, in agreement with the recent interferometric results.    Before the spectro-interferometric era, one of the first strong indications that Be disks are Keplerian came from high-resolution spectroscopy ($R>30000$) of shell stars. \\citet{han95} and \\citet{riv99} studied the so-called \\textit{Central Quasi-Emission} (hereafter, CQE) peaks, where the disk, under certain circumstances, causes a cusp in the deepest region of the line profile of shell stars. The existence of the CQE implies slow radial motions of the gas, which means the disk is supported by rotation.  In this work, we show that the same mechanism that produces  the CQE observed in shell star line profiles can cause important changes in the intensity map of the Be star plus disk system, with observable effects on the interferometric quantities. We analyze and quantify the diagnostic potential of this effect, which we dub \\textit{CQE Phase Signature} (hereafter, CQE-PS). Interestingly, similarly to the CQE detection, the CQE-PS can only be studied by high-resolution observations ($R>12000$), such as can be obtained by AMBER/VLTI in the near infrared. ",
        "conclusions": "\\label{concl} The new interferometric observations of Be stars at high spectral resolution lead to new measurements of differential phases. Deviations from the canonical appearance may occur in two different circumstances: (i) by purely interferometric properties at observations outside the marginally resolved regime, where the differential phases no longer correspond to the photocenter of the system, and (ii) by physical process accessible at high-resolution affecting the interferometric signal. In particular, in this work we discuss the role of the disk differential absorption as an important factor in shaping the differential phases of Be stars.     The solution of the radiative transfer problem of stellar photospheric flux in interaction with a typical rotating Be disk can cause a measurable photocenter displacement. Since this is an effect of differential absorption by the disk, it will depend on various system parameters and can provide valuable information about its configuration. The diagnostic potential of the CQE-PS in determining important disk parameters, such as density and velocity field, will studied in a future publication (Faes et al., in prep.)."
    },
    "1208/1208.4210_arXiv.txt": {
        "abstract": "Spectroscopic observations of metal-poor stars have indicated possible $^6$Li abundances that are much larger than the primordial abundance predicted in the standard big bang nucleosynthesis model.  Possible mechanisms of $^6$Li production in metal-poor stars include pregalactic and cosmological cosmic-ray (CR) nucleosynthesis and nucleosynthesis by flare-accelerated nuclides.  We study $^9$Be production via two-step $\\alpha$-fusion reactions of CR or flare-accelerated $^{3,4}$He through $^6$He and $^{6,7}$Li, in pregalactic structure, intergalactic medium, and stellar surfaces.  We solve transfer equations of CR or flare particles and calculate nuclear yields of $^6$He, $^{6,7}$Li, and $^9$Be taking account of probabilities of processing $^6$He and $^{6,7}$Li into $^9$Be via fusions with $\\alpha$ particles.  Yield ratios, i.e., $^9$Be/$^6$Li, are then calculated for the CR and flare nucleosynthesis models.  We suggest that the future observations of $^9$Be in metal-poor stars may find enhanced abundances originating from metal-poor CR or flare activities. ",
        "introduction": "\\label{sec1}  $^6$Li/$^7$Li isotopic ratios of metal-poor stars (MPSs) have been measured spectroscopically.  A possible plateau abundance of $^6$Li/H$\\sim 6\\times10^{-12}$ has been suggested~\\citep{asp2006}, which is about 1000 times higher than the prediction of the standard big bang nucleosynthesis (BBN) model.  Such a high abundance of $^6$Li can, however, be derived erroneously because of asymmetries in atomic line profiles originating from convective motions in atmospheres~\\citep{cay2007}.  The effect of the convection-driven line asymmetries was recently estimated and it was reported that high $^6$Li abundances have been likely detected in only a reduced number of MPSs \\citep[at most several MPSs;][]{asp2008,gar2009,ste2010,ste2012}.   The high abundance level cannot be explained in standard Galactic cosmic-ray (CR) nucleosynthesis models \\citep{men1971,ree1974,ram1997,van2000,fie2000,val2002} since the models predict $^6$Li abundances much smaller than the observed level at the low-metallicity region of [Fe/H] $<-2$~\\citep{pra2006}.  There are three different classes of astrophysical models for explanations of high $^6$Li abundances in MPSs which assume astrophysical energy sources for nuclear reactions producing $^6$Li.  The first model is the cosmological CR (CCR) nucleosynthesis model, in which $^6$Li is produced via the $\\alpha+\\alpha$ fusion and spallation of CNO nuclei~\\citep{mon1977,rol2005,rol2006,rol2011}.\\footnote{We note that Equation (1) of \\citet{rol2005}, Equation (11) of \\citet{rol2006}, and Equation (2) of \\citet{evo2008} are all wrong in the same way.  The CR injection spectrum should be given by a power law in momentum, i.e., Equation (9) of \\citet{kus2008}, in order to obtain a consistent formulation (E. Rollinde 2007, private communication).}  \\citet{evo2008} have claimed that a CCR model based on a dedicated hierarchical model of Galaxy formation does not reproduce a $^6$Li abundance level or its plateau shape.  The reason for the small $^6$Li abundance is a suppressed star formation rate at high redshift they adopted.\\footnote{We note that Equation (8) of \\citet{evo2008} for the cross section of $^4$He($\\alpha$, $X$)$^6$Li involves a typographical error.  The correct equation is $\\sigma_l(E)=66~\\exp[-0.0159(4E/{\\rm MeV})]$~mb~\\citep{mer2001}, with $\\sigma_l$ the cross section for the reaction $\\alpha+\\alpha\\rightarrow ^6$Li+$X$, and $E$ the kinetic energy per nucleon of the incident $\\alpha$ particle (C. Evoli 2012, private communication).}  Since this CCR model should include CNO spallation, $^9$Be and $^{10,11}$B are necessarily coproduced~\\citep{kus2008,rol2008}.  This model assumes CR activities at a typical redshift of $z=\\mathcal{O}$(1--10) before Galaxy formation.  The second model is the pregalactic CR (PCR) nucleosynthesis model, in which $^6$Li is produced via the $\\alpha+\\alpha$ fusion reaction between CR $\\alpha$ accelerated in structure formation shocks developed in an early epoch of Galaxy and interstellar $\\alpha$ \\citep{suz2002}.  This process operates in the epoch of structure formation until the formation of observed stars.  In the above two models, the index of the CR source spectrum with a power law in momentum should be $\\gamma \\sim 3$ for production of significant amounts of $^6$Li since the smaller and larger indexes fail to predict high $^6$Li abundances as observed in MPSs \\citep[see Figure 6 of][]{kus2008}.  The third model is nucleosynthesis by flare-accelerated energetic nuclides, in which $^6$Li is produced mainly via the $^3$He+$\\alpha$ reaction between flare-accelerated $^3$He and $\\alpha$ in stellar atmospheres \\citep{tat2007,tat2008}.  The $^3$He+$\\alpha$ reaction has been introduced in a study on $^6$Li production in solar flares \\citep{ram2000}.  Flare nucleosynthesis enhances $^6$Li in stellar atmospheres from the time of star formation to the present.  \\citet{tat2007} calculated the nucleosynthesis assuming that the source energy spectrum of flare-accelerated particles is an unbroken power law in kinetic energy of spectral index $s=4\\pm 1$~\\citep{ram1996}, and that the number ratio $^3$He/$^4$He of accelerated particles is 0.5.  Another possibility for $^6$Li production in the early universe is the $\\alpha+\\alpha$ fusion reaction~\\citep{nak2006} associated with simultaneous CNO spallation~\\citep{fie1996,fie2002,nak2004} in Type Ic supernova (SN) explosions through reactions between SN ejecta and interstellar matter (ISM) including circumstellar matter (CSM).  Type Ic SNe possibly contribute to abundances of $^6$Li, Be, and B, in the part of MPSs that are composed of material ejected from the SNe, while the three mechanisms described above somewhat homogeneously enhance the abundances.  The light element production in this model is similar to that in the PCR model except for differences in energy spectra and $^4$He abundances of projectile and target matters.  We show that this model can also produce $^9$Be via the two-step $\\alpha$-fusion reactions in this paper.  The nuclide $^9$Be is thought to be produced in the Galaxy predominantly through spallation of carbon, nitrogen, and oxygen at reactions with protons and $^4$He~\\citep[][see \\citet{pra2012} for a recent comprehensive theoretical study in light of astronomical observations]{ree1970,men1971,ree1974}.  Abundances of $^9$Be in many MPSs have been measured.  The $^9$Be abundance increases nearly in proportion to Fe abundance~\\citep{boe1999,pri2000a,pri2000b,pri2002,boe2006,tan2009,smi2009,ito2009,ric2009,boe2011}.  The severest lower limit on the primordial Be abundance, i.e., log(Be/H)$<-14$, has been deduced from an observation of the carbon-enhanced MPS BD+44$^\\circ$493 with an iron abundance [Fe/H]$=-3.7$ with Subaru/HDS \\citep{ito2009}.  The relation between abundances of Be and iron and also B~\\citep{dun1997,gar1998,pri1999,cun2000} and iron is explained by Galactic CR (GCR) nucleosynthesis models including primary and secondary reactions between CNO nuclides and $p$ and $\\alpha$~\\citep[e.g.,][]{ram1997,van2000,fie2000,val2002,pra2012}.  As for the primordial abundance of $^9$Be, a very small abundance of $^9$Be is produced in BBN at the cosmic time of $t\\lesssim 200$ s \\citep{coc2012}.  Although the $^9$Be production operates in BBN through reactions including $^6$Li($\\alpha$, $p$)$^9$Be, a thermal condition of BBN never produces observationally significant amounts of $^9$Be.  If nonthermal nuclides existed in an epoch after BBN, however, $^9$Be can be generated through $\\alpha$-fusion reactions by nonthermal $^6$He and $^6$Li nuclei which are produced via $\\alpha+\\alpha$ fusion reactions \\citep{pos2011}.  In nonthermal nucleosynthesis triggered by CRs or stellar flares, $^9$Be production by two-step nonthermal reactions also occurs.  This $^9$Be production has, however, not been taken into account in the three models for $^6$Li production in the early epoch mentioned above.  Therefore, in this paper we calculate production rates of $^9$Be as well as $^6$He and $^{6,7}$Li in the PCR, CCR, and flare nucleosynthesis models.  The structure of this paper is as follows.  The nuclear reactions that we study are explained in Section~\\ref{sec2}.  Our nucleosynthesis model for source energy spectra, nuclear transfers, reaction yields, and input parameters are described in Section~\\ref{sec3}.  Results of nuclear yields are shown in Section~\\ref{sec4}, and predicted $^9$Be/$^6$Li ratios are presented in Section \\ref{sec5}.  We summarize this study in Section~\\ref{sec6}. ",
        "conclusions": "\\label{sec6}  We have suggested the possibility of $^9$Be production by two-step $\\alpha$-fusion reactions of CR or flare-accelerated $^{3,4}$He via intermediate nuclides $^6$He and $^{6,7}$Li which occur in the IGM, pregalactic structure, and stellar surfaces.  We calculated probabilities that nuclides $^6$He and $^{6,7}$Li with variable kinetic energy at production synthesize $^9$Be through reactions with an $\\alpha$ particle.  Production rates of $^{6,7}$Li and $^9$Be are then calculated in three models, i.e., nucleosynthesis by pregalactic and cosmological CRs and flare energetic particles, by taking into account transfers and nuclear reactions of CR or flare particles.  What we have found is as follows.  1) In the pregalactic CR nucleosynthesis model, $^9$Be is produced mainly via \\\\ $^4$He($\\alpha$, $X$)$^6$Li($\\alpha$, $p$)$^9$Be.  $^6$Li has the highest probability of producing $^9$Be among all primary product nuclides.  This stems from the highest cross section of the reaction $^6$Li($\\alpha$, $p$)$^9$Be.  The calculated result of the $^9$Be/$^6$Li ratio indicates that a $^6$Li production up to the possible high level observed in MPSs is accompanied by production of $^9$Be at a level of the present observational upper limit.  An enhanced abundance of $^9$Be may, therefore, be observed in MPSs in the future if this pregalactic CR nucleosynthesis is the cause of the observed high abundance of $^6$Li.  2) In the flare nucleosynthesis model, $^9$Be is produced mainly via $^3$He($\\alpha$, $p$)$^6$Li($\\alpha$, $p$)$^9$Be.  The original seed of $^9$Be is $^3$He since the $^3$He/$^4$He ratio is assumed to be high (as in the Sun), and the softer energy spectrum prefers the reaction $^3$He($\\alpha$, $p$)$^6$Li with a lower threshold than $^4$He($\\alpha$, $X$)$^6$Li.  Since the two-step $\\alpha$-fusion reactions need rather high initial energies of seed $^{3,4}$He nuclei, the softer source spectrum in the model results in a low yield of $^9$Be relative to that of $^6$Li.  The calculated ratio $^9$Be/$^6$Li is then $\\sim 0.2$ times as large as that of the pregalactic CR model.  The $^9$Be yield in the present process is metallicity-independent lower limits on abundances in MPSs, and is larger than the yield through CNO spallation in MPSs with metallicity $Z\\lesssim 1.3\\times 10^{-2}Z_\\sun$.  3) In the cosmological CR model, $^9$Be is produced mainly via $^4$He($\\alpha$, $X$)$^6$Li($\\alpha$, $p$)$^9$Be.  Since the cosmic expansion enhances the energy loss rate of CRs, yields of nucleosynthesis decrease from those in the pregalactic CR model.  The decreases in yields of the secondary product $^9$Be are larger than those of the primary products $^{6,7}$Li since the former results from two-step energy loss processes while the latter experience energy loss only once during the CR propagation.  The calculated ratio $^9$Be/$^6$Li is then $\\sim 0.05$--$0.8$ times as large as that of the pregalactic CR model for the propagation of CRs from redshift $4\\leq z_s\\leq 30$ to 3.  When a model of CR energy profile in the universe as a function of redshift \\citep{dai2006} is adopted, the integrated ratio $^9$Be/$^6$Li is $\\sim 0.5$ times as large as that of the pregalactic CR model.  This $^9$Be yield is a metallicity-independent lower limit.  To this yield, a contribution of CNO spallation which depends on the time evolution of metallicity in the cosmic chemical evolution model should be added."
    },
    "1208/1208.0605_arXiv.txt": {
        "abstract": "We report weak-lensing masses for 51 of the most X-ray luminous galaxy clusters known. This cluster sample, introduced earlier in this series of papers, spans redshifts $0.15 \\lesssim z_{\\rm cl} \\lesssim 0.7$, and is well suited to calibrate mass proxies for current cluster cosmology experiments. Cluster masses are measured with a standard `color-cut' lensing method from three-filter photometry of each field. Additionally, for 27 cluster fields with at least five-filter photometry, we measure high-accuracy masses using a new method that exploits all information available in the photometric redshift posterior probability distributions of individual galaxies.  Using simulations based on the COSMOS-30 catalog, we demonstrate control of systematic biases in the mean mass of the sample with this method, from photometric redshift biases and associated uncertainties, to better than 3\\%. In contrast, we show that the use of single-point estimators in place of the full photometric redshift posterior distributions can lead to significant redshift-dependent biases on cluster masses. The performance of our new photometric redshift-based method allows us to calibrate `color-cut' masses for all 51 clusters in the present sample to a total systematic uncertainty of $\\approx7\\%$ on the mean mass, a level sufficient to significantly improve current cosmology constraints from galaxy clusters. Our results bode well for future cosmological studies of clusters, potentially reducing the need for exhaustive spectroscopic calibration surveys as compared to other techniques, when deep, multi-filter optical and near-IR imaging surveys are coupled with robust photometric redshift methods. ",
        "introduction": "Galaxy clusters have become a cornerstone of the experimental evidence supporting the standard $\\Lambda$CDM cosmological model.  Recent studies of statistical samples of clusters have placed precise and robust constraints on fundamental parameters, including the amplitude of the matter power spectrum, the dark energy equation of state, and departures from General Relativity on large scales.  For a review of recent progress and future prospects, see \\citet*{aem11}.  Typical galaxy cluster number count experiments require a mass-observable scaling relation to infer cluster masses from survey data, which in turn requires calibration of the mass-proxy bias and scatter. Weak lensing follow-up of clusters can be used, and to some extent has already been used, to set the absolute calibrations for the mass-observable relations employed in current X-ray and optical cluster count surveys \\citep[e.g.][]{mae08, mantz10a, vikhlinin09, rwr10}.  However, targeted weak lensing follow-up efforts of cluster surveys have not yet studied a sufficient number of clusters nor have demonstrated a sufficient control over systematic uncertainties to meaningfully impact on cosmological constraints.  For the current generation of X-ray cluster surveys drawn from ROSAT observations \\citep[e.g.][]{bcs98,reflex,fourhundred, brightmacs}, the uncertainty in the absolute mass calibration of the survey proxy, which is of the order $\\approx 15\\%$, dominates the systematic uncertainty on the matter power spectrum normalization $\\sigma_8$ \\citep{mantz10a, vikhlinin09}. For \\citet{mantz10a}, the current limits on this systematic uncertainty are derived from simulations of non-thermal pressure support in relaxed clusters \\citep[e.g.,][]{nvk07} and uncertainties in the \\textit{Chandra} calibration, whereas for \\citet{vbe09} the limits are derived from weak lensing calibrations \\citep{hoekstra07, zhang08}, quoted as a 9\\% uncertainty but neglecting an additional systematic uncertainty on the lensing masses known to be at least 10\\% \\citep{mahdavi08}. The absolute mass calibration from weak lensing follow-up therefore needs to be accurate to better than 15\\% to impact significantly on current work. Future surveys will face even more stringent systematics requirements on the absolute calibration of multiwavelength mass proxies if they are to utilize fully their statistical potential. For example, the Dark Energy Survey will require an absolute mass calibration at the 5\\% level for the dark energy constraints to be within 10\\% of their maximum potential sensitivity \\citep{heidi10}, requiring a combination of weak lensing and high-precision mass proxies, i.e. X-ray observations. Similar arguments apply to cluster surveys across the electromagnetic spectrum, e.g. the South Pole Telescope \\citep[SPT,][]{wbh11} and eRosita \\citep{pab10}.  To achieve such calibration with weak lensing, one needs to follow up a large sample of clusters. For individual clusters weak lensing typically offers mass measurements with a precision of $\\approx30\\%$ \\citep{becker11, okabe_masses, hoekstra07}, driven approximately equally by a limited number of well measured galaxies and line of sight structure.  However, simulations show that weak-lensing measurements can in principle provide accurate, approximately unbiased, estimates of the \\emph{mean} mass for statistical samples of galaxy clusters \\citep{becker11, cok07}. Small systematic biases in the mean mass can still arise from, e.g., the details of the assumed mass model, shear calibration, and the lensed-galaxy redshift distribution. Such sources of uncertainty, in particular the lensed-galaxy redshift distribution, have not yet been sufficiently understood for upcoming, or even current, surveys, as we show in this work.  In the \\emph{Weighing the Giants} project, we aim to provide absolute mass-calibration for galaxy cluster mass proxies, including specifically X-ray mass proxies, to better than 10\\% accuracy. We have gathered extensive optical imaging of 51 clusters in at least three wide photometric filters, where clusters are mostly drawn from the X-ray selected cosmological cluster sample of \\citet{mantz10a} and the relaxed cluster sample from \\citet{ars08}. Of these, 27 were observed in at least five filters. The clusters span a redshift range of $0.15 < z < 0.7$. To ensure an accurate mass-calibration, we have pursued a `blind' analysis where we have deliberately delayed comparing our lensing masses to X-ray masses and the lensing masses of others reported in the literature. Such a simple procedure prevents us from introducing observer's bias into our results. Given the redshift range, data quality, filter coverage, and blind analysis, our study represents the most extensive analysis of its type to date, and should be considered a pathfinder for the challenges facing upcoming optical, submillimetre, and X-ray cluster surveys.  Here, we report weak-lensing masses for the 51 clusters in our sample, and show that the total systematic uncertainty on the mean mass of the sample is controlled to $\\approx$7\\%. In particular, we focus on controlling systematic uncertainties associated with the redshift distribution of lensed galaxies. We approach this problem in two ways. For the entire sample, we employ a standard analysis technique \\citep[the ``color-cut'' method;][]{hoekstra07}, albeit with some improvements, where the lensed redshift distribution for each cluster field is estimated from separate, deep field photometric redshift (\\photoz) measurements. We show that this method alone does not sufficiently control systematic uncertainties to the accuracy required for current surveys. Alternatively, the redshift distribution of background galaxies may be measured using photometric redshift estimates in fields with at least five filter coverage. While previous large photometric surveys \\citep{wmk04, ilbert09} have shown that high fidelity photometric redshift point estimators are possible through the use of many (e.g., greater than 15) broad, medium, and narrow band filters for objects down to $i^{+}<25$ magnitude, observations of cluster fields usually lack coverage with such a comprehensive array of photometric filters and future optical surveys will typically have only six broad filters. We show that with such limited photometric coverage, \\photoz point estimates are insufficient to recover unbiased cluster masses. We therefore develop a method that uses the full \\photoz posterior probability distribution \\pz for individual galaxies in each cluster field, referred to as the ``$P\\left ( z\\right )$'' method, and show that it can be used to measure robust cluster weak-lensing masses. Using the COSMOS-30 \\photoz catalog \\citep{ilbert09}, we create a series of simulations to test the sensitivity of our reconstructed masses to \\photoz errors. We show that $P\\left ( z\\right )$ distributions from current photometric redshift codes, with \\BVRiz photometry, enable control of systematic uncertainties on the mean mass for the sample to better than 2\\% accuracy for clusters at $0.15 < z < 0.7$ -- a result that provides significant encouragement for future cluster-cosmology work.  \\label{contents}  This is the third in a series of papers describing the project. Paper I describes our cluster sample, data reduction procedures, and shear measurements \\citep{paper1}. Paper II  details our photometric redshift measurements, including the development of a scattered-light correction for SuprimeCam, and an improved relative photometric calibration procedure based on fitting the stellar locus \\citep{paper2}. This paper reports our lensing masses and estimates of systematic uncertainties in the sample mean mass. Forthcoming papers will use these accurate cluster masses to calibrate X-ray mass proxies and determine improved cosmological constraints.   The structure of this paper is as follows. In Section~\\ref{sec:lensing_theory}, we review cluster mass measurements with weak lensing. We describe our data set and analysis procedures in Section~\\ref{sec:data_reduction}. In Section~\\ref{sec:color_cuts}, we develop and apply our implementation of the color-cut method to all clusters in the sample. In Sections~\\ref{sec:maxlike_method} \\& \\ref{sec:shape_distribution}, we introduce our \\photoz lensing framework that incorporates \\photoz posterior probability distributions for each galaxy observed.  In Section~\\ref{sec:testing_framework}, we investigate the expected systematic errors present in mass measurements given the empirical performance of \\photoz estimators. In Section~\\ref{sec:final_masses}, we report measured masses using both the \\pz and color-cut methods, and cross-calibrate the color-cut results. In Section~\\ref{sec:systematics}, we perform checks of other potential systematic uncertainties. In Section~\\ref{sec:aveconcentration}, we briefly digress to measure the average concentration of massive clusters. We compare our lensing mass measurements to other efforts in the literature, based on overlapping samples, in Section~\\ref{sec:lit_comp}.  Finally, we provide concluding remarks in Section~\\ref{sec:conclusions}.   Unless otherwise noted, all mass measurements assume a flat $\\Lambda CDM$ reference cosmology with $\\Omega_{\\rm m} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ and $H_0 = 100 \\,h\\, \\mbox{km/s/Mpc}$, where $h=0.7$. ",
        "conclusions": "Systematic uncertainties in our mass calibration arise from the shear measurements, the assumed mass model, and from uncertainties in the lensed-galaxy redshifts. We have investigated each in turn with checks from simulations and data. Table~\\ref{table:syserr_summary} summarizes our estimates of the systematic uncertainties in the analysis.  The systematic uncertainties associated with the shear measurements and the mass model apply \\emph{both} to the color-cut and the \\pz methods. The uncertainties associated with redshift measurements listed in the table apply only to the \\pz method, and are constrained to be less than 2\\%.  The systematic uncertainties that only apply to the color-cut method are more difficult to quantify (Section~\\ref{sec:color_cut_sys_errs}). To gauge the uncertainties in those measurements, we therefore pursue a strategy of cross-calibration. By scaling the color-cut masses by the average ratio between the two methods (as measured in Section~\\ref{sec:final_masses}), we calibrate out the unknown systematic uncertainties in the color-cut analysis for the price of adding the statistical uncertainty in the average ratio. For the 27 clusters with \\BVRiz photometry, the uncertainty in the cross-calibration between the color-cut and \\pz methods is $\\approx 4\\%$.  \\begin{table*} \\caption{Summary of the sources and levels of systematic uncertainty in the analysis. Shear measurement and mass model uncertainties apply equally to the color-cut and \\pz methods. The quoted redshift measurement uncertainties apply only to the \\pz method; the color-cut method is subject to other, more difficult to quantify systematic uncertainties discussed in the text when not cross-calibrated from the \\pz method. Additionally, systematic uncertainties are quoted for all 51 clusters studied with the color-cut method, and 27 clusters that use the \\pz masses directly.  Values in the table are reported to single-digit precision.} \\begin{tabular}{l | c c c  } \\hline Uncertainty Source & \\multicolumn{3}{c}{\\% of Mean Cluster Mass} \\\\ \\hline & Color-Cut Method & & P(z) Method \\\\ \\hline \\textbf{Shear Measurements} & &  &\\\\ Multiplicative Shear Bias Cor & & 3\\% & \\\\ STEP PSF Mismatch & & 2\\% & \\\\ Coaddition \\& PSF Interpolation & & 1\\% & \\\\ \\hline \\textbf{Mass Model} &  & & \\\\ Profile Uncertainty & & 3\\% & \\\\ \\hline \\textbf{Photo-$z$ Measurements} & &  & \\\\ Residual Photometry Systematics & & 3\\% & \\\\ Simulated Photo-$z$ Bias & &1\\% & \\\\ Depth \\& Filter Mismatch & &1\\% & \\\\ \\hline \\textbf{Method Cross-Calibration}& 4\\% & & -  \\\\ \\hline \\hline \\textbf{Method Systematic Uncertainty} & 7\\% &  & 6\\%  \\\\ \\hline Triaxiality \\& LOS Structure & 3\\% &  & 4\\% \\\\ \\hline \\hline \\textbf{Total Systematic Uncertainty} & 8\\% & & 7\\% \\\\  \\end{tabular} \\label{table:syserr_summary} \\end{table*}  We expect each source of systematic uncertainty to be independent, and have approximated each source as a Gaussian. Our total systematic uncertainty on the mean cluster mass, for 51 clusters, is therefore 7\\%. Results are comparable when only masses measured with the \\pz method are used."
    },
    "1208/1208.5020_arXiv.txt": {
        "abstract": "{The contribution of the merging process to the early phase of galaxy assembly at $z > 1$ and, in particular, to the build-up of the red sequence, still needs to be accurately assessed.} {We aim to measure the major merger rate of star-forming galaxies at $0.9 < z <1.8$, using close pairs identified from integral field spectroscopy (IFS).} {We use the velocity field maps obtained with SINFONI/VLT on the MASSIV sample, selected from the star-forming population in the VVDS. We identify physical pairs of galaxies from the measurement of the relative velocity and the projected separation ($r_{\\rm p}$) of the galaxies in the pair. Using the well constrained selection function of the MASSIV sample, we derive at a mean redshift up to $z = 1.54$ the gas-rich major merger fraction (luminosity ratio $\\mu = L_2/L_1 \\geq 1/4$), and the gas-rich major merger rate using merger time scales from cosmological simulations.} {We find a high gas-rich major merger fraction of $20.8^{+15.2}_{-6.8}$\\%, $20.1^{+8.0}_{-5.1}$\\%, and $22.0^{+13.7}_{-7.3}$\\% for close pairs with $r_{\\rm p} \\leq 20h^{-1}$ kpc in redshift ranges $z = [0.94, 1.06], [1.2, 1.5)$, and $[1.5, 1.8) $, respectively. This translates into a gas-rich major merger rate of $0.116^{+0.084}_{-0.038}$~Gyr$^{-1}$, $0.147^{+0.058}_{-0.037}$~Gyr$^{-1}$, and $0.127^{+0.079}_{-0.042}$~Gyr$^{-1}$ at $z = 1.03, 1.32$, and $1.54$, respectively. Combining our results with previous studies at $z < 1$, the gas-rich major merger rate evolves as $(1+z)^{n}$, with $n = 3.95 \\pm 0.12$, up to $z = 1.5$. From these results we infer that $\\sim35$\\% of the star-forming galaxies with stellar masses $\\overline{M}_{\\star} = 10^{10}-10^{10.5}\\ M_{\\odot}$ have undergone a major merger since $z \\sim 1.5$. We develop a simple model that shows that, assuming that all gas-rich major mergers lead to early-type galaxies, the combined effect of gas-rich and dry mergers is able to explain most of the evolution in the number density of massive early-type galaxies since $z \\sim 1.5$, with our measured gas-rich merger rate accounting for about two-thirds of this evolution.} {Merging of star-forming galaxies is frequent at around the peak in star formation activity. Our results show that gas-rich mergers make an important contribution to the growth of massive galaxies since $z \\sim 1.5$, particularly on the build-up of the red sequence.} ",
        "introduction": "Understanding the mechanisms involved in the mass assembly of galaxies and their relative role over cosmic time is an important open topic in modern astrophysics. In particular, the evolution of the red sequence, which includes passive galaxies dominated by old stellar populations and an early-type (E/S0) morphology, imposes fundamental constraints on the formation and evolution models.  The stellar mass density in the red sequence has increased by a factor of $\\sim 10$ in the 2.5 Gyr between $z = 2$ and $z = 1$, but only by a factor of $\\sim 2$ in the last $7 - 8$ Gyr of cosmic history \\citep[e.g.,][]{arnouts07,vergani08,ilbert10}. Major mergers, the merger of two galaxies with similar stellar masses, is an efficient mechanism for creating new passive, early-type galaxies \\citep[e.g.,][]{naab06ss,rothberg06a,rothberg06b,hopkins08ss,rothberg10,bournaud11}. Thus, the knowledge of the merger rate at $z > 1$ is important input when estimating the relative contribution of merging and cold-gas accretion \\citep[e.g.,][]{dekel06} in the early assembly of galaxies and, in particular, the role of merging in the build-up of the red sequence.  The evolution of the merger rate since $z \\sim 1$ is now well constrained by direct observations. The early measurements using photometric pairs \\citep{patton97,lefevre00} or post-merger morphological signatures \\citep{conselice03,jogee09} have been superseded by spectroscopic measurements confirming physical pairs from the redshift measurement of both components of a major merger with a luminosity/mass ratio $\\mu \\geq 1/4$ \\citep[e.g.,][]{lin08,deravel09,deravel11}, as well as for minor mergers down to $\\mu = 1/10$ \\citep{clsj11mmvvds}. With a parametrization of the merger rate's evolution following $\\propto (1 + z)^{n}$, it is observed that the major merger rate's evolution depends on the luminosity and on the mass of the galaxy sample \\citep[e.g.,][]{deravel09}, where massive galaxies with $M_{\\star} > 10^{11}\\ M_{\\odot}$ have a higher merger rate, but with little redshift evolution ($n \\sim 0 - 2$), while lower mass galaxies with $M_{\\star} = 10^{9}-10^{11}\\ M_{\\odot}$ have a lower merging rate but with stronger redshift evolution ($n \\sim 3 - 4$). This mass dependency seems to explain some of the apparent discrepancy of merger rate measurements made from observations targeting different mass samples.  Beyond $z \\sim 1$, direct measurements of the merger rate are still limited. Previous attempts to measure the major merger rate at $z > 1$ have focused on the identification of merger remnants from morphological signatures \\citep{conselice08,conselice11,bluck12}, on the study of projected close pairs \\citep{ryan08,bluck09,williams11,man12,marmol12,law12}, or on indirect estimations \\citep{cameron12,puech12}. These studies find a high merger rate to $z\\sim2-3$ but with a large scatter between different measurements. However, these results are up to now solely based on photometric measurements which are increasingly hard to correct for contamination along the line of sight as redshift increases. Another complication stems from the morphological evolution of galaxies, with show more irregular morphologies at high redshifts, and a wavelength dependency with more multi-component objects present when observed in the rest-frame UV \\citep{law07}, with some of these components possibly related to strong star-forming regions rather than to different dynamical components.  To improve on this situation, it is necessary to obtain spectroscopic confirmation of the physical nature of the photometric pairs at $z \\gtrsim 1$. In the last years, NIR Integral Field Spectrographs (IFSs), like SINFONI on the VLT or OSIRIS on the Keck, have opened the possibility for a systematic study of the dynamical field around high redshift galaxies in the optical rest-frame. Some examples are the MASSIV\\footnote{http://www.ast.obs-mip.fr/users/contini/MASSIV/} (Mass Assembly Survey with SINFONI in VVDS, \\citealt{massiv1}) survey at $0.9 < z < 1.8$, the SINS\\footnote{http://www.mpe.mpg.de/~forster/SINS/sins\\_nmfs.html} (Spectroscopic Imaging survey in the Near-infrared with SINFONI, \\citealt{sins}) survey at $z \\sim 2$, or the Keck-OSIRIS \\citep{keckosiris}, the AMAZE (Assessing the Mass-Abundance redshift -Z- Evolution, \\citealt{amaze}) and the LSD (Lyman-break galaxies Stellar populations and Dynamics, \\citealt{lsd}) surveys at $z \\sim 3$.  The MASSIV survey has been designed to target the peak of the star-formation rate at $0.9 < z < 1.8$, filling the gap between higher redshift ($z \\sim 2$) IFS surveys with those at $z < 1$, e.g., IMAGES (Intermediate MAss Galaxy Evolution Sequence, \\citealt{images}). The MASSIV survey has targeted 84 star-forming galaxies at $0.9 < z < 1.8$ with SINFONI , drawn from the VVDS\\footnote{http://cesam.oamp.fr/vvdsproject/} (VIMOS VLT Deep Survey, \\citealt{lefevre05}) survey. MASSIV has been used as a unique opportunity to study in detail the dynamical state of $0.9 < z < 1.8$ galaxies \\citep{massiv2}, their metallicity gradients \\citep{massiv3}, or the evolution of the fundamental mass-size-velocity relations since $z \\sim 1.2$ \\citep{massiv4}.  In this paper, using the MASSIV survey, we present for the fist time a measurement of the gas-rich major merger rate of star-forming galaxies from kinematical close pairs at $0.9 < z < 1.8$. Thanks to the large field-of-view of IFS we have access to the complete surrounding volume of the galaxies when searching for close kinematical companions. In addition, the well-defined selection of sources from the VVDS and the well controlled selection function of MASSIV observations ensures the study of a representative population of star-forming galaxies at these redshifts \\citep[see][for details]{massiv1}. This all together enables the measurement of average volume quantities like the merger fraction and rate.  The paper is organised as follows: in Sect.~\\ref{data} we summarise the MASSIV data set used to identify merging pairs, and in Sect.~\\ref{method} we develop the methodology to measure the merger fraction from IFS data. We report the gas-rich major merger fraction in MASSIV in Sect.~\\ref{ffmassiv}, and derive the gas-rich major merger rate in Sect.~\\ref{mrmassiv}. We discuss the implication of our results in Sect.~\\ref{discussion}. Finally, we present our conclusions in Sect.~\\ref{conclusion}. We use $H_0 = 100h$ km s$^{-1}$ Mpc$^{-1}$, $h = 0.7$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_{\\Lambda} = 0.7$ throughout this paper. All magnitudes refer to the AB system. The stellar masses assume a \\citet{salpeter55} initial mass function (IMF).   \\begin{figure*}[t] \\centering \\includegraphics[width = 9cm]{wsel_wide14.eps} \\includegraphics[width = 9cm]{wsel_wide22.eps} \\includegraphics[width = 9cm]{wsel_deep.eps} \\includegraphics[width = 9cm]{wsel_udeep.eps} \\caption{MASSIV selection in the [$\\ion{O}{ii}$]$\\lambda3727$ flux [10$^{-17}$ erg s$^{-1}$ cm$^{-2}$] vs equivalent width [$\\AA$] plane (see \\citealt{massiv1}, for details). Red squares are the VVDS sources at $0.94 < z < 1.5$ with individual $[\\ion{O}{ii}$]$\\lambda3727$ line measurement in the VVDS-Wide 14h (top-left), VVDS-Wide 22h (top-right), VVDS-Deep (bottom-left), and VVDS-Ultradeep surveys (bottom-right). The dashed lines mark the selection of MASSIV star-forming galaxies. White dots are those VVDS sources that fulfil the MASSIV selection. Green pentagons are the MASSIV galaxies observed with SINFONI/VLT. [{\\it A colour version of this plot is available at the electronic edition}].} \\label{wmassiv} \\end{figure*} ",
        "conclusions": "Using SINFONI/VLT 3D spectroscopy, we have been able to measure, for the first time with spectroscopically-confirmed close pairs, the gas-rich major merger fraction and merger rate at around the peak in star formation activity at $0.9 < z < 1.8$, from the MASSIV sample of star-forming galaxies with a stellar mass range $M_{\\star} = 10^{9} - 10^{11}\\ M_{\\odot}$. In this redshift range we identify 20 close pairs, and classify 13 as major mergers based on a separation $r_{\\rm p} \\leq 20h^{-1} - 30h^{-1}$ kpc and a relative velocity $\\Delta v \\leq 500$ km~s$^{-1}$.  We find that the gas-rich major merger fraction is high, $20.8^{+15.2}_{-6.8}$\\%, $20.1^{+8.0}_{-5.1}$\\%, and $22.0^{+13.7}_{-7.3}$\\% for $r_{\\rm p} \\leq 20h^{-1}$ kpc close pairs in redshift ranges $z = [0.94, 1.06], [1.2, 1.5)$, and $[1.5, 1.8)$, respectively. When compared to measurements at redshifts $z < 1$, the evolution of the (gas-rich) merger fraction can be parametrised as $f_{\\rm MM} = 0.0066 \\times (1 + z)^{m}$ with $m = 3.91 \\pm 0.16$. We note that the evolution between $z = 1$ and $z \\sim 1.5$ seems to flatten out compared to lower redshifts.  The merger rate has been derived using merger time scales from the literature: we find that the gas-rich merger rate is $0.116^{+0.084}_{-0.038}$~Gyr$^{-1}$, $0.147^{+0.058}_{-0.037}$~Gyr$^{-1}$, and $0.127^{+0.079}_{-0.042}$~Gyr$^{-1}$ at $z = 1.03, 1.32$, and $1.54$, respectively, for merger time scales of $T_{\\rm MM} \\sim 1.5$ Gyr. We then find that the (gas-rich) merger rate's evolution for galaxies with stellar mass $\\overline{M}_{\\star} = 10^{10-10.5}\\ M_{\\odot}$ over $z = [0, 1.5]$ scales as $(1 + z)^n$ with $n = 3.95 \\pm 0.12$.  Using these measurements, we developed a simple model to estimate the contribution of gas-rich major mergers to the growth of galaxies on the red sequence. We infer that $\\sim35$\\% of the star-forming galaxies with stellar mass $\\overline{M}_{\\star} = 10^{10}-10^{10.5}\\ M_{\\odot}$ have undergone a major merger since $z \\sim 1.5$. The number of major merger events was about 5 times higher at $1 < z < 1.5$ compared to $z < 1$. Assuming that each gas-rich major merger produces a new early-type galaxy, we infer that the number of gas-rich major mergers is large enough at $z > 1$ to explain the increase in the number density of massive ETGs, supporting a picture where gas-rich (wet) merging is the main process building-up the red sequence. While gas-rich mergers become rarer towards lower redshifts, the number of dry mergers is steadily increasing, and the combination of these two processes accounts for most, if not all, of the increase in the number density of massive ETGs since $z \\sim 1.3$. Two-thirds of this number density evolution is due to wet major mergers, while one-third is coming from major and minor dry mergers. These results add further evidence to a picture where merging is a major process driving mass assembly into the massive red sequence galaxies.  We note that minor merging is definitely present in the MASSIV sample (see Sect.~\\ref{ffmassiv}). However, due to incompleteness in detecting these faint companions, we are not able to assess a minor merger rate at these epochs from our data. In the global picture of red sequence assembly, we emphasise that a simple extrapolation of the minor merger rate measured up to $z \\sim 1$ by \\citet{clsj11mmvvds}, would indicate that from $z \\sim 1.5$ to the present, minor mergers with $1/10 \\leq \\mu < 1/4$ are not common enough to significantly move spiral galaxies into the red sequence.  To get a complete picture of the life-time effect of major merging on massive galaxies, accurate measurements of the merger fraction and merger rate are needed beyond $z \\sim 2$. Spectroscopic surveys will remain an important element to provide secure identification of merging systems at these early epochs."
    },
    "1208/1208.5166_arXiv.txt": {
        "abstract": "\\noindent We show that ``accidental'' supersymmetry is a beyond-the-Standard Model framework that naturally accommodates a  thermal relic dark matter candidate and successful electroweak baryogenesis, including the needed strongly first-order character of the electroweak phase transition. We study the phenomenology of this setup from the standpoint of both dark matter and baryogenesis.  For energies around the electroweak phase transition temperature, the low-energy effective theory is similar to the MSSM with light super-partners of the third-generation quarks and of the Higgs and gauge bosons.  We calculate the dark matter relic abundance and the baryon asymmetry across the accidental supersymmetry parameter space, including resonant and non-resonant CP-violating sources. We find that there are regions of parameter space producing both the observed value of the baryon asymmetry and a dark matter candidate with the correct relic density and conforming to present-day constraints from dark matter searches. This scenario makes sharp predictions for the particle spectrum, predicting a lightest neutralino mass between 200 and 500 GeV, with all charginos and neutralinos within less than a factor 2 of the lightest neutralino mass and the heavy Higgs sector within 20-25\\% of that mass, making it an interesting target for collider searches. In addition, we demonstrate that successful accidental supersymmetric dark matter and baryogenesis will be conclusively tested with improvements smaller than one order of magnitude to the current performance of electron electric dipole moment searches and of direct dark matter searches, as well as with IceCube plus Deep Core neutrino telescope data. ",
        "introduction": "The Standard Model (SM) of particle physics is missing several key ingredients needed for a satisfactory phenomenological description of nature.  First, it does not provide an explanation for the observed baryon asymmetry of the universe (BAU). Second, the SM does not contain any viable particle candidates for dark matter (DM), which is needed to explain a large array of astrophysical and cosmological observations.  From a more theoretical perspective, the SM additionally falls short of explaining the large hierarchies between fundamental physical scales.  In particular, it provides no satisfactory explanation for why the Planck scale, $M_{\\rm Pl}\\sim 10^{19}$ GeV, is so much higher than the electroweak (EW) scale, $m_{\\rm EW}\\sim 100$ GeV.  This is known as the hierarchy problem.  In recent years, several models have been suggested that address the hierarchy problem.  In particular, models with warped extra dimensions, such as Randall-Sundrum (RS) scenarios, have been proposed which naturally generate the hierarchy between the Planck and EW scales \\cite{Randall:1999ee}.  In RS models, the universe is described by a five-dimensional (5D) geometry with two four-dimensional (4D) branes located at the UV (Planck) and IR (TeV) points.  The Higgs fields are localized on (or near) the IR brane, and the warped fifth dimension ``redshifts'' the Planck scale to the TeV scale, providing a rather elegant solution to the hierarchy problem.  Additionally, by placing the SM fermions in the bulk, hierarchies between the Yukawa couplings can be accounted for by the wave-function overlap with the Higgs boson in the fifth dimension \\cite{Gherghetta:2000qt}.  While explaining the hierarchy problem, simply embedding the SM in a RS scenario is not fully satisfactory.  To prevent sizable CP-violating effects from Kaluza-Klein (KK) modes in the absence of additional flavor structure, the IR scale must be at or above $\\mathcal{O}(10$ TeV$)$\\cite{KK_modes}.  Precision electroweak experiments also dictate that the IR scale must be larger than the EW scale, hence some additional tuning is required between these scales.  This is an incarnation of the so-called \\emph{little} hierarchy problem.  To resolve this issue, models of ``emergent\" or ``accidental\" supersymmetry have been proposed (see e.g. Refs.~\\cite{Gherghetta:2003wm, Sundrum:2009gv, accidental}), in which supersymmetry (SUSY) emerges as an accidental symmetry in the IR, with SUSY broken on the UV brane.  As a result, the Higgs mass can be protected from radiative corrections up to the IR scale, while the warped extra dimension generates the hierarchy between the TeV and Planck scales.  Within this framework, which we describe in more detail in Sec.~\\ref{sec:Accidental_SUSY}, both hierarchy problems can potentially be resolved.  The specific particle content of the theory depends on the model of SUSY embedded in the Randall-Sundrum spacetime.  Since we are interested in the general features of accidental supersymmetric models, we will consider the particle content of the minimal supersymmetric extension of the standard model (MSSM) as a conservative case from the standpoint of the field content of the theory.  Randall-Sundrum scenarios can do more than just solve the big and little hierarchy problems.  In fact, we show here that models with warped extra dimensions may also provide an explanation for the origin of the baryon asymmetry via the mechanism of electroweak baryogenesis (EWB). Crucial to this mechanism are a number of conditions, ultimately related to those generically needed for any dynamical mechanism for the production of a baryon asymmetry \\cite{Sakharov:1967dj}: first, one needs departure from thermal equilibrium at the electroweak scale; second, one needs large enough charge (C) and charge-parity (CP) violation; third, one needs violation of baryon number. The second and third conditions are easily satisfied in the context of supersymmetric EWB models: any minimal supersymmetric extension to the SM, in fact, contains numerous (albeit constrained) new sources of CP violation, while baryon number (B) violation, is provided by SM weak sphalerons --- we will comment on this more below. More critical is how to have a large deviation from thermal equilibrium --- a condition in practice realized, in the context of EWB, via a strongly first-order electroweak phase transition. As pointed out in Ref.~\\cite{Nardini:2007me}, here the RS setup may be of crucial importance, as we also explain below.  While the universe is described by the Randall-Sundrum spacetime at zero temperature, at finite temperature, RS models possess an additional high-temperature phase, described by an Anti-de Sitter-Schwarzschild (AdS-S) spacetime with a black hole horizon replacing the TeV brane \\cite{Creminelli:2001th}.  Alternatively, a holographic description facilitated by the AdS-CFT correspondence also exists in which the two phases correspond to a deconfined and to a confined phase of a strongly coupled gauge theory, respectively.  Provided that the free energy of the RS phase is less than the free energy of the AdS-S phase, $F_{RS}<F_{AdS-S}$, a phase transition can occur between the two: as the universe cools below a temperature $T_c$, bubbles of the TeV brane can begin to nucleate out of the black hole horizon \\cite{Creminelli:2001th} (see also Refs.~\\cite{Randall:2006py, Kaplan:2006yi, Nardini:2007me} for further discussion of the confining phase transition).  Because, from the CFT perspective, conformal invariance is only spontaneously broken in the RS phase, $F_{RS}>F_{AdS-S}$ implies that the RS phase is metastable.  For the confining phase transition to occur, one must introduce some mechanism to explicitly break conformal invariance.  From the AdS perspective, this can be accomplished by stabilizing the radion (the field governing the separation between the UV and IR branes) with a potential generated e.g. by additional 5D fields.  Once the radion is stabilized, the free energy of the two phases will be equal at some temperature $T=T_c$ producing a phase transition via bubble nucleation with nucleation temperature $T_n\\leq T_c$.  In many cases, $T_n$ can be significantly lower than the temperature of the electroweak phase transition (EWPT) predicted by the 4D Minkowski theory \\cite{Nardini:2007me}.  Since the Higgs sector is typically confined to the IR brane, this low nucleation temperature results in a ``{\\em supercooled}'' EWPT (i.e. taking place at lower temperatures than otherwise possible), thereby potentially strengthening the phase transition.  While this supercooling was studied specifically in the case of the SM embedded in RS with a Goldberger-Wise potential \\cite{Goldberger:1999uk} for the radion in Ref.~\\cite{Nardini:2007me}, this possibility is a consequence of the geometry and localization of the Higgs sector in the IR and is largely independent of the particle content of the theory and can therefore potentially arise in accidental SUSY as well.  As a result, models of accidental SUSY may provide a strongly first order EWPT even \\emph{without} e.g. a light right-handed scalar top (stop) quark \\cite{Carena:2008vj, Carena:2008rt}, or additional singlets contributing to the Higgs potential  \\cite{Profumo:2007wc}, as is typically required for successful EWB in the MSSM.  Alternatively, as we explain in the next section, certain incarnations of the accidental SUSY framework also posit, as a solution to the $\\mu$-problem, an additional extension to the Higgs sector via a singlet scalar field. This potentially provides an additional route to a strongly first order electroweak phase transition.  Since a strongly first order phase transition appears to be a natural possibility in accidental SUSY, the remaining issue pertinent to EWB is the requirement of large enough CP-violation to seed weak sphalerons, which will be the primary focus for the rest of this study.  In fact, accidental SUSY naturally satisfies this requirement as well.  Even in its MSSM incarnation, there are several new CP-violating phases which can source the baryon asymmetry. In particular, there are new phases arising in the higgsino-gaugino and third-generation scalar sectors.  It has recently been shown that of the third-generation scalars, only CP-violating stau sources can account for the observed baryon asymmetry while still conforming to various phenomenological and experimental constraints from electric dipole moment searches \\cite{Kozaczuk:2012xv}.  Here we concern ourselves with moderate values of the ratio of Higgs vevs, $\\tan\\beta =10$, in which case the stau sources are suppressed.  We will therefore be interested in EWB with higgsino-gaugino sources in accidental SUSY.  Electroweak baryogenesis utilizing these sources in the MSSM has been extensively analyzed in recent studies (e.g. \\cite{Huet:1995sh, Carena:1996wj, Lee:2004we, Chung:2008aya, Lepton_Mediated, Supergauge, Including_Yukawa, Konstandin:2003dx, Konstandin:2004gy, More_Relaxed, Carena:2002ss, Carena:2008vj, EWB_and_EDMs, Balazs:2004ae, EWB_and_DM, Kozaczuk:2011vr, Menon:2004wv, Huber:2006wf}), and we build on these analyses in our study of the accidental SUSY scenario.    Note that extending the particle content beyond that of the MSSM would provide more potential  sources of CP-violation.  Supersymmetric RS models also have the added benefit of generically containing a viable dark matter candidate, if the lightest supersymmetric particle (LSP) corresponds to the lightest neutralino, over some regions of parameter space.  This is a result of $R$-parity conservation, whereby the LSP is stable.  Thus it may be possible for accidental SUSY to simultaneously explain the origin of the BAU and the nature of dark matter, while also solving both the big and little hierarchy problems.  In fact, the production of both the relic DM density and the baryon asymmetry via higgsino-gaugino sources are closely connected \\cite{EWB_and_DM, Kozaczuk:2011vr}, since both depend predominantly on the higgsino mass term $\\mu$ and on the gaugino soft-supersymmetry breaking masses $M_1$ and $M_2$.  Consequently, enforcing both the correct DM properties and baryon asymmetry in conformity with various observational constraints may result in sharp predictions for the regions of interest within the accidental SUSY parameter space.  This is an attractive possibility and one which we explore in the present study.  In what follows we compute the baryon asymmetry across the parameter space of a minimal (MSSM-like) incarnation of accidental SUSY.  We do so independently of the specifics of SUSY breaking, choosing higgsino and gaugino masses which yield the correct DM relic density (this is the so-called ``well-tempered neutralino\" setup \\cite{welltempered}) and assuming a strongly first-order electroweak phase transition arising either from the supercooling provided by the AdS-S transition or from the contribution of a gauge singlet super-field to the effective potential.  We then impose constraints from electric dipole moment (EDM) measurements and from dark matter searches to outline potentially viable regions of the parameter space.  We also consider the impact of the projected sensitivities of these various experiments.  In doing so, we find that accidental SUSY models, even in their most minimal incarnations, may allow for successful EWB and a viable DM candidate provided that the resulting soft breaking wino mass, $M_2$, and the higgsino mass parameter $\\mu$ are roughly degenerate, with the LSP a bino-higgsino admixture; we also require the MSSM heavy Higgs sector to lie relatively close to twice the LSP mass.  These findings can be used to hone in on specific models of accidental SUSY giving rise to the observed properties of our universe. ",
        "conclusions": "\\label{sec:conclusions} Accidental supersymmetry is a particle physics framework that naturally addresses both the large and the little hierarchy problems as well as the potential CP and flavor problems of supersymmetry, while in principle providing a successful thermal dark matter candidate. We argued here that this framework naturally accommodates successful electroweak baryogenesis, for the following reasons: \\begin{itemize} \\item[(i)] a strongly first order electroweak phase transition may be a generic feature of this framework, either as a consequence of supercooling produced by the phase transition between the high and low-temperature RS spacetimes, or from the contribution of a singlet to the superpotential as may be required to solve the $\\mu$-problem; \\item[(ii)] light third-generation stops and gauginos allow for resonant CP-violating sources to produce potentially large net chiral currents fueling a large enough net baryon number via sphaleron transitions; \\item[(iii)] heavy first- and second-generation sfermions prevent excessive one-loop contributions to observable electric dipole moments in the presence of the needed large CP-violating phases. \\end{itemize} Here, we carried out a model-independent study (from the standpoint of supersymmetry breaking), although for definiteness we picked a specific accidental SUSY spectrum realization. Specifically, we let the relevant U(1)$_Y$ and SU(2) gaugino soft supersymmetry breaking masses $M_1$ and $M_2$, as well as the higgsino mass parameter $\\mu$ vary freely. We constrained this triplet of mass parameters enforcing that the lightest supersymmetric particle be a neutralino with a thermal relic density matching the observed density of dark matter. In practice, this amounted to selecting values of $\\mu$ across the $(M_1,\\ M_2)$ parameter space so that the higgsino fraction drove the thermal relic density of the lightest neutralino to the desired value.  After enforcing the relic density constraint, we proceeded to calculate the baryon asymmetry resulting from electroweak baryogenesis across the $(M_1,\\ M_2)$ parameter space. We included both resonant and non-resonant sources, and we picked two representative values for the heavy Higgs sector mass scale, which is relevant for resonant sources. The requirement of successful baryogenesis generically restricted the viable parameter space to a relatively narrow funnel at $M_1\\lesssim M_2$, with $\\mu\\sim M_1,\\ M_2$; recent direct detection constraints also enforce $M_1\\simeq m_A/2$ to within 20-25\\%.  The strongest constraints on this framework derive from the non-observation of electric dipole moments and of signals from dark matter direct detection, most notably with the Xenon100 experiment \\cite{xenon100}. We calculated in detail how these constraints restrict the parameter space relevant for baryogenesis, concluding that dark matter direct searches eliminate neutralinos with a large higgsino fraction (requiring to some degree resonant annihilation through the heavy Higgs sector, and hence $M_1\\simeq m_A/2$), while electric dipole moments greatly restrict regions of viable electroweak baryogenesis to those parameter space points producing, for maximal CP violating phases, a BAU much larger than observed (those parameter space regions are then compatible with successful baryogenesis as the CP phases are lowered to comply with EDM searches).  We calculated the predicted EDM and dark matter search rates in the framework of accidental supersymmetric baryogenesis and we concluded that: \\begin{itemize} \\item the most sensitive EDM search to constrain this model is provided by searches for the electron EDM; an improvement of one order of magnitude on the current experimental sensitivity would conclusively test the framework, even allowing for some theoretical uncertainties in the calculation of the BAU; \\item the entire parameter is highly constrained by current direct, spin-independent dark matter-nucleon cross section limits, and will soon be fully tested even for resonant neutralino annihilation \\item the predicted signal at neutrino telescopes from neutralino annihilation in the core of the Sun is potentially large enough for detection, although direct detection results imply that no signal is expected within approximately one year of data taking \\end{itemize}  The parameter space compatible with successful baryogenesis and thermal dark matter is highly constrained, and is characterized by a lightest neutralino with a mass between 200 and 500 GeV, with all other neutralino and chargino masses within a factor 2 of the lightest neutralino mass, with $m_A\\simeq 2M_1<1$ TeV. This compressed electroweak ``inos'' spectrum might be challenging for LHC searches, but would be ideally targeted with an $e^+ e^-$ linear collider with a TeV center of mass energy.  Concluding, we demonstrated here that accidental supersymmetry is an explicit realization of a framework for successful thermal relic dark matter and electroweak baryogenesis, which is motivated by an entirely different set of theoretical arguments based upon addressing the hierarchy, CP and flavor problems. We showed that accidental supersymmetric baryogenesis is a highly constrained setup, but one with very sharp experimental predictions for electric dipole moment, dark matter, and collider searches. We therefore anticipate that this scenario be falsified or produce signals in the very near future in a variety of experiments."
    },
    "1208/1208.5485_arXiv.txt": {
        "abstract": "Using a novel technique, we achieve ${\\sim}$100 picoarcsecond resolution and set an upper bound of less than 4 km for the characteristic size of the Vela pulsar's emission region. Specifically, we analyze flux-density statistics of the Vela pulsar at 760 MHz. Because the pulsar exhibits strong diffractive scintillation, these statistics convey information about the spatial extent of the radio emission region. We measure both a characteristic size of the emission region and the emission sizes for individual pulses. Our results imply that the radio emission altitude for the Vela pulsar at this frequency is less than 340 km. ",
        "introduction": "\\label{sec::Intro} \\subsection{Pulsar Emission and Scintillation}  Owing to their extraordinary emission, pulsars are invaluable probes of a vast array of extreme physics, from the supranuclear density interiors and the ${\\sim}10^{12}\\mathrm{\\ G}$ magnetospheres of neutron stars, to the turbulent plasma of the interstellar medium (ISM). The regularity of their time-averaged emission has been applied to sensitive tests of general relativity \\citep{Taylor_Weisberg_89,Kramer_GR}, and the achieved timing precision holds promise for direct observation of gravitational waves \\citep{Detweiler_GW,Jenet_GW,IPTA}. Yet, remarkably, no consensus exists for the precise origin of this emission; its geometry, stability, and physical mechanism remain enigmatic, even after nearly a half-century of study \\citep{Melrose_2004,Verbiest_2009}. The incredible compactness of the pulsar and its surrounding plasma prohibit traditional imaging of the emission process, but a number of methods can achieve the requisite resolution, using the interstellar plasma as an effective `lens' \\citep{Lovelace_Thesis,Backer_75,Cordes83,CornwellNarayan93,ISO,Shishov_2010}.  Variations of this stochastic lens give rise to the observed scintillation, and statistics of these variations convey information about the spatial structure of the emission. Longer observing wavelengths enhance the signature of this structure because the stronger scattering effectively enlarges the aperture. However, the increased spatial resolution comes at the expense of spectral resolution, and it becomes difficult to decouple the scintillation from source noise and variability. Indeed, the pulse-broadening timescale can easily exceed that of intrinsic pulsar variation.  We can circumvent these difficulties by constructing spectra over accumulation times that include \\emph{all} pulsed power. Intrinsic sub-pulse variability then affects the correlation of spectral noise but not the mean spectrum or single-channel statistics \\citep{Intermittent_Noiselike_Emission}. Furthermore, we can distinguish the effects of extended source geometry from the scattering evolution and pulsar variability by varying the degree of temporal averaging; with no averaging, only the source emission structure contributes. Additionally, we can identify the emission sizes of individual pulses. We have presented a detailed mathematical treatment separately \\citep[hereafter JG12]{IPDF}.  We constructed spectra in this way and then compared histograms formed from the spectral samples with appropriate models for a strongly scintillating source. Such models benefit from minimal assumptions about the nature and distribution of both the scattering material and the underlying emission; this generality is achieved at the expense of resolution in pulse phase. The inferred emission size is therefore a characteristic size of the emission region, quantified via a dimensionless parameter $\\gamma_{\\mathrm{s}}$, which depends on the spatial standard deviation of source intensity, weighted by integrated flux density. Translating $\\gamma_{\\mathrm{s}}$ to physical units at the pulsar additionally requires specification of the scattering geometry.    \\subsection{Comparison with Previous Work} \\label{sec::Previous} Several previous investigations have also quantified scintillation statistics in order to estimate the size of the Vela pulsar's emission region. In general, for a strongly scattered point source, the scintillation acts upon intensity as a multiplicative gain $\\mathcal{G}$ that is the squared modulus of a complex random walk. The probability density function (PDF) of $\\mathcal{G}$ is therefore exponential, $P(\\mathcal{G}) = e^{-\\mathcal{G}}$, with an associated modulation index of unity: $m^2 \\equiv \\left\\langle \\mathcal{G}^2 \\right \\rangle / \\left\\langle \\mathcal{G} \\right \\rangle^2 - 1 = 1$. A spatially-extended emission region smoothes this variation, and so the modified $P(\\mathcal{G})$ or reduced modulation index can be used to estimate the emission size \\citep{Salpeter67,Cohen67,ISO}.  The principal difficulty in detecting these modifications is the additional noise that arises from the noiselike nature of the source, known as \\emph{self-noise}, and from the background. Pulsars are particularly challenging targets because their dramatic intrapulse and interpulse variations also contribute modulation. All of these sources of noise can easily imitate or mask the signature of an extended emission region. Previous efforts have attempted to minimize these sources of noise by spectral averaging in both frequency and time. In contrast, the present technique fully accounts for the effects of self-noise and intrinsic pulsar variability.  For example, \\citet{Gwinn97} sought to detect the modification of $P(\\mathcal{G})$ through the PDF of the correlation function amplitude on the Tidbinbilla-Parkes baseline at 2.3 GHz. Because this baseline was much shorter than the scale of the diffraction pattern, the observational setup was effectively a zero-baseline interferometer. Their best-fitting model corresponded to an emission region with a FWHM of $460 \\pm 110$ km. \\citet{Gwinn00} then extended this analysis to three gates across the pulse, also at 2.3 GHz. They improved the model PDF and quantified the effects of their averaging in frequency and time. They estimated a declining size of the emission region from $440\\pm 90\\mathrm{\\ km}$ to less than $200 \\mathrm{\\ km}$ across the pulse.  Rather than analyzing distribution functions, \\citet{Macquart00} used the modulation index to quantify the effects of emission size for single-dish data. They measured $m^2 \\approx 0.9$ at 660 MHz and thereby inferred the extent of source emission to be no more than 50 km ($\\mathrm{FWHM}=120$ km). Including the contribution of self-noise leads to an inferred FWHM of the emission region of ${\\sim}200\\mathrm{\\ km}$ \\citepalias{IPDF}.   More recently, \\citet{Vela_18cm} analyzed visibilities on the Tidbinbilla-Mopra baseline at 1.65 GHz and fit models to the projections of the real part and imaginary variance. They found the size of the emission region to be large both early and late in the pulse (${\\sim}400\\mathrm{\\ km}$ and ${\\sim}800\\mathrm{\\ km}$, respectively) but near zero in the central portion, where the pulse is strong.   An analysis at lower frequencies has the advantage of increasing the signature of size but the disadvantage of incurring more mixing between the scintillation and the pulsar noise. We demonstrate that this difficulty is not prohibitive, and that a careful statistical description of observed power spectra at low frequencies can yield fresh insight into the physics of both interstellar scattering and pulsar emission.     \\subsection{Outline of Paper}  In \\S\\ref{sec::Theory}, we briefly review the theoretical descriptions of pulsar emission and interstellar scintillation. Then, in \\S\\ref{sec::Obs}, we describe our observations and the subsequent construction of both the observed and model PDFs of intensity; we also present the expected residual structure arising from an extended emission region and outline the model fitting procedure. Next, in \\S\\ref{sec::Analysis}, we give our analysis for various pulse selection cuts and derive our estimates for the size of the emission region. In \\S\\ref{sec::Errors_and_Bias}, we quantify possible biases and systematic errors in our inferences of the size of the emission region. Then, in \\S\\ref{sec::Interpretation}, we interpret our measurements in terms of beamed, dipolar emission, and derive the corresponding emission altitudes. Finally, in \\S\\ref{sec::Summary}, we summarize our results, their impact on our understanding of pulsar magnetospheres, and the prospects for future work. ",
        "conclusions": "\\label{sec::Summary}  We have analyzed flux-density statistics of the Vela pulsar at 760 MHz. In particular, we examined the observed PDF of intensity in various cuts determined by frequency, polarization, and pulse strength. Based on single pulse estimates of the background noise and gated signal-to-noise, we constructed model PDFs in order to identify the signature arising from a spatially-extended emission region.  The predicted form for point-source emission showed excellent agreement with the data ($\\chi^2_{\\mathrm{red}} {\\approx} 1$), good to ${\\sim}0.01\\%$ over the low to moderate intensity regions of the PDF, in which the effects of an extended emission region are the most pronounced. This type of comparison requires no fitted parameters and arises from minimal assumptions about the characteristics of both the scattering (i.e.\\ strong and slowly evolving) and source emission (i.e.\\ amplitude-modulated noise with small spatial extent). Our inferences are therefore extremely robust.  We did not detect any effects arising from an extended emission region. We achieved the tightest limit by jointly analyzing all pulses in three 3.2 MHz subbands, centered on 751.25, 757.5, and 763.75 MHz; this limit corresponds to a characteristic size of the emission region of $\\sigma_{\\mathrm{c}} \\ltwid 4\\mathrm{\\ km}$ at the pulsar at $3\\sigma$ confidence. This limit involved information about the nature and distribution of the scattering material. If the phase structure function of the scattering material is Kolmogorov, for instance, then the limit decreases to $\\sigma_{\\mathrm{c}} \\ltwid 2\\mathrm{\\ km}$.  The inferred size, $\\sigma_{\\mathrm{c}}$, reflects the spatial extent of all the emitting regions during a pulse and, therefore, also gives an upper-bound on the instantaneous emission size, which we assume is pointlike in our derivation of the emission altitude. Explicitly, $\\sigma_{\\rm c}$ corresponds to the standard deviation of a circular Gaussian distribution of intensity that yields equivalent scintillation statistics as those observed. Conversion to an emission altitude also involves the effective angular pulse width, weighted by flux density, which provides a natural mapping to lateral emission structure. We thereby obtained constraints on the emission altitude: $r_{\\mathrm{em}} < 340\\mathrm{\\ km}$ (square-law) or $r_{\\mathrm{em}} < 170\\mathrm{\\ km}$ (Kolmogorov).  We also individually tested pulses for the signature of an extended emission region, again without a statistically significant detection. For the strongest pulses we obtained limits $\\ltwid 50\\mathrm{\\ km}$, and for the majority of pulses we obtained limits $\\ltwid 200 \\mathrm{\\ km}$. However, even for the strongest pulses, this bound only marginally constrained the emission to be generated within the light cylinder. Nevertheless, our results demonstrate the application and reliability of this technique.  Our upper bound on the size of the emission region improves previous estimates using scintillation statistics by over an order of magnitude. Without invoking symmetry constraints or assumptions about the distribution of emission within the open field-line region, we achieve estimates of the emission altitude that are comparable with alternative techniques that analyze profile widths and polarization characteristics. Thus, even for incredibly compact `core' emission, scintillation statistics now have the ability to refine and extend the information conveyed by bulk polarimetry.  The excellent agreement between our data and models demonstrates that our technique can be similarly applied to other pulsars without fear of spurious inference. In particular, the quality of our data appears to have controlled the systematic uncertainties to be within the fundamental restriction set by Poisson noise. Thus, future observations, perhaps at lower frequencies or of other pulsars with favorable scattering geometries, may confidently measure a spatially-extended emission region and thereby greatly enrich our empirical understanding of pulsar radio emission."
    },
    "1208/1208.2093_arXiv.txt": {
        "abstract": "We present the first spectroscopic study of the globular clusters (GCs) in the giant elliptical galaxy (gE) M86 in the Virgo cluster. Using spectra obtained in the Multi-Object Spectroscopy (MOS) mode of the Faint Object Camera and Spectrograph (FOCAS) on the Subaru telescope, we measure the radial velocities for 25 GCs in M86. The mean velocity of the GCs is derived to be $\\overline{v_p}=-354^{+81}_{-79}$ \\kmss, which is different from the velocity of the M86 nucleus ($v_{\\rm gal}=-234\\pm 41$ \\kmss). We estimate the velocity dispersion of the GCs, $\\sigma_p=292^{+32}_{-32}$ \\kmss, and find a hint of rotation of the M86 GC system. A comparison of the observed velocity dispersion profiles of the GCs and stars with a prediction based on the stellar mass profile strongly suggests the existence of an extended dark matter halo in M86. We also estimate the metallicities and ages for 16 and 8 GCs, respectively. The metallicities of M86 GCs are in the range $-2.0<$ [Fe/H] $<-0.2$ with a mean value of $-1.13 \\pm 0.47$. These GCs show a wide age distribution from 4 to 15 Gyr. ",
        "introduction": "Globular clusters (GCs) are excellent tracers for studying the formation history of their host galaxies \\citep{lee03, bro06}. In particular, a giant elliptical galaxy (gE) contains thousands of GCs, located close to the galaxy center to very far away in the outer halo. Therefore, GCs can be used as powerful test particles for studying the kinematics and chemical evolution of gE halos.  Up to now, there are several kinematic studies of the GC systems in nearby gEs: M49 \\citep{zep00,cot03}, M60 \\citep{pie06, bri06,lee08b, hwa08}, M87 \\citep{coh97,kis98b,cot01}, NGC 4636 \\citep{sch06, cha08, par10, lee10a}, NGC 1399 \\citep{kis98, min98, kis99, ric04, ric08, sch10}, NGC 5128 \\citep{pen04a, pen04b, woo07}, and NGC 1407 \\citep{rom09}. The data used for these studies were compiled and reanalyzed in \\citet{lee10a}, who % found that the kinematic properties of the GC systems are diverse among the gEs,  indicating diverse merging and accretion histories of gEs (see also the recent study on the M87 GC system \\citep{str11}).    There are also several studies focusing on the spectroscopic ages and metallicities of the gE GCs: \\citet{coh98} for M87, \\citet{bea00} and \\citet{coh03} for M49, \\citet{pen04b}, \\citet{bea08}, \\citet{woo10}, and \\citet{woo10b} for NGC 5128, \\citet{pie06} for M60, \\citet{cen07} for NGC 1407, and \\citet{kis98} and \\citet{for01} for NGC 1399. Recently, \\citet{par12a} presented a study of spectroscopic ages and metallicities for the GCs in NGC 4636. They also compiled the data of all the gE GCs in the literature, and found that the GC metallicity distribution in the combined gE sample is bimodal.  We have been carrying out a project to investigate the spectroscopic properties of GCs in nearby galaxies to understand the formation of the GC systems in galaxies. % Our study on the kinematics of the GC system of the Virgo gE M60 was presented in \\citet{lee08b} and \\citet{hwa08} and that on the GC system of the spiral galaxy M31 % in \\citet{kim07} and \\citet{lee08c}. Recently, we presented measurements of the radial velocities for the GCs in NGC 4636 \\citep{par10}, % and a detailed kinematic analysis of these data in \\citet{lee10a}. We also investigated the chemical properties of NGC 4636 GCs and % other gE GCs \\citep{par12a}.  Here, we present a spectroscopic study of the GCs in M86 (NGC 4406), a gE in the Virgo cluster. This galaxy is one of the best targets for the spectroscopic study of the GC system because it is located close to the center of the Virgo cluster and harbors GCs out to large radii from the galaxy center \\citep{lee10b}. To date, there has been no published spectroscopic study of M86 GCs.  In contrast to the absence of spectroscopic studies, there have been several studies of the photometric properties of M86 GCs. For example, \\citet{kun01} and \\citet{lar01} found that the color distribution of M86 GCs is  bimodal from an analysis of HST/ WFPC2 images. % This bimodality was confirmed by HST/ACS data \\citep{pen06} and ground-based wide-field imaging data \\citep{rho04,par12b}. The radial number density profile of the M86 GC system is approximately well fitted by a de Vaucouleurs law and power law \\citep{rho04,par12b}. Basic information about M86 is summarized in Table \\ref{tab-m86gal}.    This paper is composed as follows. Section 2 describes the spectroscopic target selection, observation, and data reduction. In \\S3, we identify genuine M86 GCs and list the corresponding photometric and spectroscopic data. We present the kinematic properties of  the M86 GC system in \\S4, and the metallicities and ages of the M86 GCs in \\S5. The primary results are summarized in the final section. ",
        "conclusions": "Using the Subaru spectroscopic data of M86 GCs, we studied the kinematic and chemical properties of the M86 GC system. Our main results are summarized as follows.  \\begin{enumerate}  \\item For the first time, we measured the radial velocities of 31 objects in the M86 field: 25 M86 GCs, % two foreground stars, one probable intracluster GC in the Virgo cluster, two faint galaxies, and the M86 nucleus.  \\item The mean velocity of the GCs is $\\overline{v_p}=-354^{+81}_{-79}$ \\kmss, which is different from the velocity of the M86 nucleus ($v_{\\rm gal}=-234\\pm 41$ \\kmss). The velocity dispersion of the GCs is $\\sigma_p=292^{+32}_{-32}$ \\kmss. The M86 GC system shows a hint of rotation.  \\item From a comparison of the VDPs predicted from the stellar mass profile with the observed VDPs of the stars and GCs, we found evidence for the existence of an extended dark matter halo in M86.  \\item We determined the metallicities for 16 GCs using the BH method, and the ages and metallicities for 8 GCs using the grid method. The metallicity of the M86 GCs derived from the BH method is in the range $-2.0<$ [Fe/H] $<-0.2$  with a mean value of $-1.13 \\pm 0.47$. The grid method results in similar [Fe/H] values and a mean age of 9.7 $\\pm$ 4.0 Gyr.    \\end{enumerate}"
    },
    "1208/1208.4095_arXiv.txt": {
        "abstract": "I describe the last years of the 370-year long life of the Sterrekundig Instituut Utrecht, which was the second-oldest university observatory in the world and was closed in early 2012 after the Faculty of Science and the Board of Utrecht University decided, without providing qualitative or quantitative arguments, to remove astrophysics from its research and education portfolio. ",
        "introduction": "\\subsection{Difficult Start for the New Faculty of Science}  In November 2007, Alfred Bliek became the new dean of the Science Faculty of Utrecht University, a dean who turned out to be a strong supporter of astronomy.  The merging of different departments into the new Science Faculty a few years earlier had led to financial difficulties because some departments brought very substantial, structural deficits into the new faculty, thereby impacting other, well-funded departments. Despite being financially sound and having substantial reserves before the merge, the Physics \\& Astronomy department had gotten under financial pressure and was forced to lay off some staff. No SIU staff was laid off. The reserves had been blocked by the university board and directly or indirectly used to cover the deficits of other departments. During these insecure times, Norbert Langer, then director of SIU, won a prestigious Alexander von Humboldt Professorship and departed at the end of 2008 for Bonn, Germany.  As of 1.1.2009, the Faculty of Science had appointent me as Scientific Director of SIU. Despite the word 'scientific' in the title, most of the actual work had to do with administration, bureaucracy and politics.   \\subsection{A New Beginning}  In April 2009 the Sterrekundig Instituut Utrecht (SIU) embarked on an ambitious plan, with full support from the dean. An additional staff position was approved to reduce the heavy teaching load at the BSc level. The position was filled by Maureen van den Berg in June 2010. The MSc program was enhanced with an instrumental track and more space-related topics to attract students from technical universities, which was reflected in its expanded name: Astrophysics and Space Research. The plans for a new, common building with SRON, NOVA (Nederlandse Onderzoekschool voor de Astronomie, the federation of Dutch universities with astronomy departments) and the faculty\u2019s Scientific Instrumentation workshop started taking shape in the summer of 2009. The new building would have allowed an effective sharing of expensive machinery and cleanrooms between all partners. In the fall of 2009, the Board of Utrecht University approved additional support for collaborations with SRON (Netherlands Institute for Space Research), providing SIU and IMAU (Institute for Marine and Atmospheric research Utrecht) with 340,000 Euros per year for 5 years. In January 2010 the NOVA board expressed its plan to move its Optical/IR group (10-15 FTE) to Utrecht around 2017/2018 to realize the vision of Dutch astronomy to concentrate all its optical/infrared instrumentation efforts in one place on a university campus. SIU and NOVA were also ready to implement their plan to replace Norbert Langer with a new full professor and two assistant professors, which was presented to the faculty in October 2010. Most importantly, all these plans were achievable with level funding from the university.   \\subsection{Top Astronomy in The Netherlands}  Astronomy in the Netherlands is a top-scoring science. The 2010 evaluation of leading research schools reaffirmed this: NOVA and Zernike (material sciences in Groningen) were labeled exemplary, meaning above excellent. Indeed, Dutch astronomy is comparable in performance to the top US institutes. This research excellence is supported by excellence in education. SIU\u2019s research was assessed as nationally leading in its areas of expertise, internationally competitive and making a significant contribution to the field. The review committee recommended recruiting a world-leading researcher in an exciting field of astrophysics and one or two additional faculty hires at more junior levels, in agreement with SIU\u2019s plans. In terms of the faculty\u2019s support of SIU\u2019s ambition, the International Review Board (IRB) noted: {\\em \u201cThe IRB was very impressed with the clear vision of the Utrecht University administration regarding the situation, and hopes that the healthy collaborative atmosphere will lead to a satisfactory resolution to the challenges facing the astrophysics group.\u201d} ",
        "conclusions": ""
    },
    "1208/1208.4740_arXiv.txt": {
        "abstract": "Until now, investigating the early stages of galaxy formation has been primarily the realm of theoretical modeling and computer simulations, which require many physical ingredients and are challenging to test observationally.  However, the latest Hubble Space Telescope observations in the near infrared are shedding new light on the properties of galaxies within the first billion years after the Big Bang, including our recent discovery of the most distant proto-cluster of galaxies at redshift $z\\sim 8$. Here, I compare predictions from models of primordial and metal-enriched star formation during the dark ages with the latest Hubble observations of galaxies during the epoch of reionization. I focus in particular on the luminosity function and on galaxy clustering as measured from our Hubble Space Telescope Brightest of Reionizing Galaxies (BoRG) survey. BoRG has the largest area coverage to find luminous and rare $z\\sim8$ sources that are among the first galaxies to have formed in the Universe. ",
        "introduction": "The first billion years after the Big Bang represent a key area of astrophysics, with interesting problems, open questions and potential for unexpected and unusual discoveries as highlighted by the 2010 Astronomy Decadal Survey Report\\footnote{\\url{http://www.nap.edu/catalog.php?record_id=12951}}. Significant progress has been made in the field in the last few years. Numerical simulations are following formation of the first stars at progressively high resolution (see reviews by Yoshida and Abel in this volume) and simulations of first galaxies from first principles are now growing in dynamic range \\citep{romanodiaz11,wise11}. Theoretical models can be constructed to investigate and predict the star formation rate during and after the epoch of Reionization at $z\\gtrsim 4$ \\citep{ts09,trenti10} (see Fig.~\\ref{fig:sfr}).  Observationally, the quest for direct detection of first stars seems very difficult, with these sources being too faint (a $100~M_{\\odot}$ Pop III star at $z>6$ has observed magnitude $M_{AB}\\gtrsim 38$). However, prospects of indirect detection of metal-free stars are intriguing, either through high-$z$ supernovae, or if gravitational lensed and living in small clusters \\citep{tss09,zack12}.  In terms of observations of high-redshift galaxies, recent progress has been done thanks to the installation of Hubble's WFC3, leading to the identification of large samples of galaxies at $z\\sim 7-10$ (see review by Bouwens in this volume).  In this contribution, I briefly summarize some of my recent results on star and galaxy formation at high redshift, starting from a short discussion of metal-free versus metal enriched star formation during the epoch of reionization (and spatial inhomogeneities in chemical enrichment), and then focusing on observations of the most overdense regions in the universe at $z\\sim8$ through the Brightest of Reionizing Galaxies survey. ",
        "conclusions": ""
    },
    "1208/1208.1812_arXiv.txt": {
        "abstract": "The prompt emission spectra of gamma-ray bursts (GRBs) usually have a dominant component that is well described by a phenomenological ``Band'' function. The physical origin of this spectral component is debated. Although the traditional interpretation is synchrotron radiation of non-thermal electrons accelerated in internal shocks or magnetic dissipation regions, a growing trend in the community is to interpret this component as modified thermal emission from a dissipative photosphere of a GRB fireball. We analyze the time dependent spectrum of GRB 110721A detected by {\\em Fermi} GBM and LAT, and pay special attention to the rapid evolution of the peak energy $E_p$. We define a ``death line'' of thermally-dominated dissipative photospheric emission in the $E_p - L$ plane, and show that $E_p$ of GRB 110721A at the earliest epoch has a very high $E_p \\sim 15$ MeV that is beyond the ``death line''. Together with the finding that an additional ``shoulder'' component exists in this burst that is consistent with a photospheric origin, we suggest that at least for some bursts, the ``Band'' component is not from a dissipative photosphere, but must invoke a non-thermal origin (e.g. synchrotron or inverse Compton) in the optically thin region of a GRB outflow. We also suggest that the rapid ``hard-to-soft'' spectral evolution is consistent with the quick discharge of magnetic energy in a magnetically-dominated outflow in the optically thin region. ",
        "introduction": "\\label{sec:intro}  The prompt emission spectrum of a gamma-ray burst (GRB) is usually well described by a phenomenological function known as the ``Band'' function (Band et al. 1993). This model, which is essentially a broken power law function with a smooth (exponential) transition, was traditionally invoked to model spectra of GRBs detected by BATSE on board the Compton Gamma-Ray Observatory. The function is found successful to describe most GRB spectra detected by later missions as long as the spectral band is wide enough (e.g. Abdo et al. 2009; Zhang et al. 2011).  The physical origin of this phenomenological Band function is not identified. The traditional model is synchrotron emission of non-thermal electrons in an optically thin region, e.g. internal shocks or internal magnetic dissipation regions (M\\'esz\\'aros et al. 1994; Tavani 1996; Daigne \\& Mochkovitch 1998; Lloyd \\& Petrosian 2000; Bosnjak et al. 2009; Zhang \\& Yan 2011). Alternatively, a matter-dominated outflow (fireball) can have a bright photosphere (Pacz\\'ynski 1986; Goodman 1986; M\\'esz\\'aros \\& Rees 2000;  M\\'esz\\'aros et al. 2002), which may be enhanced by kinetic or magnetic dissipation processes near the photosphere (Thompson 1994; Rees \\& M\\'esz\\'aros 2005; Pe'er et al. 2006; Giannios 2008; Beloborodov 2010; Lazzati \\& Begelman 2010; Ioka 2010). It has been argued that due to geometrical and/or physical broadening, this quasi-thermal component may be modified to mimic a Band function (e.g. Beloborodov 2010; Lazzati \\& Begelman 2010; Pe'er \\& Ryde 2011; Lundman et al. 2012). The scalings of a non-dissipative photosphere model are argued to be able to interpret various empirical correlations (Fan et al. 2012).  Within the framework of the standard fireball-shock model, a GRB prompt emission spectrum is expected to be the superposition of a quasi-thermal photosphere emission component and a non-thermal component in the optically-thin internal shock region (M\\'esz\\'aros \\& Rees 2000; Zhang \\& M\\'esz\\'aros 2002; Toma et al. 2011; Pe'er et al. 2012). Such a superposition effect has been claimed in the BATSE data archive (e.g. Ryde 2005; Ryde \\& Pe'er 2009), and was confirmed more robustly recently with the Fermi data (e.g. Ryde et al. 2010; Zhang et al. 2011; Guiriec et al. 2011; Axelsson et al. 2012). On the other hand, most GRB spectra (e.g. GRB 080916C) are still well described by one single Band component (Abdo et al. 2009; Zhang et al. 2011). This sharpens the debate regarding the origin of the Band function. For GRB 080916C, the non-detection of a thermal component led to the suggestion of a Poynting-flux-dominated outflow (Zhang \\& Pe'er 2009; see also Daigne \\& Mochkovitch 2002; Zhang \\& M\\'esz\\'aros 2002). Alternatively, some authors attempted to interpret the entire Band function as emission from a dissipative photosphere (e.g. Beloborodov 2010; Vurm et al. 2011; Giannios 2008; Ioka 2010). These two interpretations invoke distinct assumptions regarding the composition of the GRB jets. Finding observational clues to differentiate between them is therefore essential to unveil the physics of GRB central engine, jet composition and energy dissipation mechanisms, which are poorly constrained (e.g. Zhang 2011).  Here we show that the time-resolved spectral information of GRB 110721A holds the key to address this open question. ",
        "conclusions": "\\label{sec:summary}  We have shown that the Band function component of GRB 110721A is beyond the ``death line'' of thermally-dominated dissipative photosphere models in the $E_p - L$ plane. Together with the fact that an additional shoulder thermal component is consistent with the photosphere model, we reach the conclusion that the so-called ``Band'' component is {\\em not} of a photospheric origin at least for this burst, and is formed via non-thermal dissipation processes in the optically thin regions. Veres et al. (2012) interpreted this emission as synchrotron emission from a magnetically-dominated jet.  Such a finding has profound implications in understanding the origin of other Band function spectra of GRBs. Zhang et al. (2011) identified three elemental spectral components through a detailed time-resolved spectral analysis of 17 GRBs co-detected with Fermi GBM and LAT. They found that there are two types of GRBs. GRB 080916C's spectra remain the ``Band'' shape even though the time interval for the spectral analysis progressively reduces, reaching $\\sim 1$ s in the rest frame. GRB 090902B, on the other hand, showed a clear ``narrowing'' feature as time bin reduces, and spectrum is much narrower at $\\sim$ rest-frame 1 s. The time integrated spectrum of this burst is also narrower than other Band GRBs. This burst's ``Band'' component can be indeed de-composed as superposed photosphere emission (Ryde et al. 2010; Zhang et al. 2011; Pe'er et al. 2012; Mizuta et al. 2011). However, such bursts are not common. Most bursts are similar to GRB 080916C.  The spectral parameters of GRB 110721A are similar to those of GRB 080916C and many other GRBs. The conclusion that the Band component of GRB 110721A originates from the optically thin region also supports the suggestion that most Band components are non-thermal (synchrotron or SSC) emission in the optically thin region. Available data seem to suggest the following unified picture: The GRB central engine may have a range of magnetization parameter $\\sigma_0$. 1. For low $\\sigma_0$ bursts, dissipation can occur at small radii, so that a bright photospheric emission component (such as the case of GRB 090902B) is detected. Such bursts are usually accompanied by a high energy component due to upscattering of the thermal photons (and probably also synchrotron self-Compton, Pe'er et al. 2012). 2. For intermediate $\\sigma_0$ bursts, the photosphere component is weaker but still detectable. Examples of this category include GRB 110721A (Axelsson et al. 2012) and GRB 100724B (Guiriec et al. 2011). The Band component of these bursts are formed in the large radii via internal shocks (IS) or internal collision-induced magnetic reconnection and turbulence (ICMART). 3. Finally, if $\\sigma_0$ is large enough, both the photosphere and IS components are suppressed. The Band component forms at even larger radii via the ICMART process (Zhang \\& Yan 2011).  Finally, very rapid hard-to-soft $E_p$ evolution observed in GRB 110721A (Axelsson et al. 2012 and this work) and many other GRBs (Lu et al. 2010, 2012) challenges existing models. Such a rapid evolution is not expected in the internal shock model. For the photosphere model, some moderate hard-to-soft evolution may be expected due to the intial growth of optical depth in a fireball (W. Deng \\& B. Zhang 2012, in preparation), but never an extreme evolution like GRB 110721A. A possible interpretion may be made within the framework of the ICMART model. According to this model (Zhang \\& Yan 2011), $\\sigma$ in the emission region rapidly decreases during each ICMART event, since the magnetic energy is continuously dissipated. If a good fraction of local magnetic dissipation energy is deposited to electrons, the typical electron Lorentz factor $\\gamma_e$ would have a $\\sigma$-dependence, so that $E_p$ decreases with time as $\\sigma$ reduces. Also relativistic turbulent reconnection would lead to locally Doppler-boosted mini-jets, whose Lorentz factors would also depend on the local $\\sigma$ value (Zhang \\& Zhang 2012). The time dependent Doppler boosts for mini-jets  would enhance the hard-to-soft $E_p$ evolution."
    },
    "1208/1208.3026_arXiv.txt": {
        "abstract": "If dark matter has a large self-interaction scattering cross section, then interactions among dark-matter particles will drive galaxy and cluster halos to become spherical in their centers.   Work in the past has used this effect to rule out velocity-independent, elastic cross sections larger than $\\sigma/m \\simeq 0.02\\hbox{ cm}^2/\\hbox{g}$ based on comparisons to the shapes of galaxy cluster lensing potentials and X-ray isophotes. In this paper, we use cosmological simulations to show that these constraints were off by more than an order of magnitude because (a) they did not properly account for the fact that the observed ellipticity gets contributions from the triaxial mass distribution outside the core set by scatterings, (b) the scatter in axis ratios is large and (c) the core region retains more of its triaxial nature than estimated before. Including these effects properly shows that the same observations now allow dark matter self-interaction cross sections at least as large as $\\sigma/m = 0.1\\hbox{ cm}^2/\\hbox{g}$.  We show that constraints on self-interacting dark matter from strong-lensing clusters are likely to improve significantly in the near future, but possibly more via central densities and core sizes than halo shapes. ",
        "introduction": "The nature of dark matter is one of the most compelling mysteries of our time.  On large scales, the behavior of dark matter is consistent with what cosmologists of yore called ``dust'' \\citep[e.g.,][]{tolman1934}, meaning its behavior is consistent with being collisionless and non-relativistic (``cold'') for the vast majority of the Universe's history \\citep{breid2010}.  This consistency has been of great interest to the particle-physics community because the most popular candidate for dark matter, the supersymmetric neutralino, displays exactly this behavior \\citep{steigman1985,griest1988b,jungman1996}.  While the supersymmetric neutralino paradigm is attractive in many ways, there are two outstanding problems with it.  First, astroparticle searches have yet to turn up evidence for the existence of the neutralino, though searches are rapidly increasing their sensitivity to interesting neutralino parameter space \\citep{geringer-sameth2011,ackermann2011a,cotta2011,fox2011b,bertone2012,atlas2012,koay2012,baudis2012,baer2012,xenon2012}.  Second, there are predictions for the structure of dark-matter halos that have not been observationally verified at a quantitative level \\citep{dubinski1991,diemand2008,vogelsberger2009,stadel2009,navarro2010}.  In fact, there are hints of tension with the neutralino paradigm on sub-galactic scales \\citep{dobler2006,gentile2007,deblok2008,kuzio2008,kuzio2010,zwaan2010,boylan-kolchin2011,papastergis2011}. And yet, most of the effort to characterize the evolution of the Universe (experimentally, observationally, and theoretically) has been in the context of this dust-like cold dark matter (CDM).  In the absence of evidence for \\emph{any} dark-matter candidate, much less the neutralino, it is important to explore the structure and evolution of the Universe for dark-matter phenomenology beyond cold and collisionless.   One intriguing possibility is that the dark matter belongs to and interacts with a ``dark'' or ``hidden'' sector \\citep{khlopov2006,feldman2007,foot2007,pospelov2008,feng2008,arkanihamed2009,feng2009,sigurdson2009,cohen2010}.  The Standard Model has rich intra-sector phenomenology; it is not unreasonable to speculate that a complex dark sector only tenuously connected to the Standard Model might exist as well.  The simplest kind of dark-matter interaction is a hard-sphere interaction of identical dark-matter particles. Such an interaction -- with an isotropic, velocity-independent, elastic scattering cross section -- was first introduced in an astrophysical context by \\citet{spergel2000}.  Interactions of this type were invoked to ameliorate the tensions between observations and CDM predictions on small scales (on the scales of individual dark-matter halos) while leaving the large-scale successes of CDM intact. In this paper we revisit this basic class of self-interacting dark matter (SIDM) using cosmological simulations to explore its effect on dark matter halo shapes as a function of cross section.  In a companion paper (Rocha et al. 2012) we investigate implications for dark matter halo substructure and density profiles.   We are reinvestigating this simple SIDM model, which had been decreed ``uninteresting'' in several studies a decade ago, for two primary reasons.  First, we suspected that the constraints that indicated that the SIDM cross section was too small to meaningfully alter the morphology of dark-matter halos, were not as tight as claimed.  Second, there is a wealth of new data (e.g., from near-field cosmology, lensing studies of galaxies and clusters) that may be better places to either look for SIDM or constrain its properties.  In this paper and our companion paper, Rocha et al. (2012), we reevaluate past constraints on SIDM and suggest several new places to look for the effects of SIDM on halo structure.  There was a burst of work on SIDM before its untimely demise \\citep{yoshida2000,yoshida:weak:2000,kochanek2000,hogan2000,dalcanton2001,dave2001,colin2002,hennawi2002}.  The death of isotropic, velocity-independent, elastically scattering SIDM largely came from the interpretation of three types of observations: halo evaporation in galaxy clusters \\citep{gnedin2001}, cores in galaxy clusters \\citep{yoshida:weak:2000,meneghetti2001}, and halo shapes \\citep{miralda2002}.  The Y2K-era constraints from the former two classes of observations were at the level of $\\sigma/m \\lesssim 0.3\\hbox{ cm}^2/\\hbox{g}$ and $\\sigma/m\\lesssim 0.1\\hbox{ cm}^2/\\hbox{g}$ respectively.  However, we show in a companion paper, Rocha et al. (2012), that the evaporation and cluster-core constraints are likely overestimated.  The most stringent constraints on dark matter models with large isotropic, elastic self-scattering cross sections emerged from the shapes of dark-matter halos, in particular from lens modeling of the galaxy cluster MS 2137-23 by \\citet{miralda2002}.  This massive galaxy cluster has a number of radial and tangential arcs within $\\sim 200$ kpc of the halo center \\citep{mellier1993,miralda1995}.  \\citet{miralda2002} argued that self interactions should make dark-matter halos round within the radius $r$ where the local per-particle scattering rate equals the Hubble rate $\\Gamma(r) = \\, H_0$, or equivalently, where each dark-matter particle experiences one interaction per Hubble time.   The scattering rate per particle as a function of $r$ in a halo scales in proportion to the local density and velocity dispersion, \\begin{eqnarray}\\label{eq:gamma} \\Gamma(r) \\sim \\rho(r) (\\sigma/m) v_{\\mathrm{rms}}(r) \\, , \\end{eqnarray} where $\\rho$ is the local dark-matter mass density and $v_{\\mathrm{rms}}$ is the rms speed of dark-matter particles.    Using the fact that the lens model needs to be elliptical at $70 \\hbox{ kpc}$,  \\citet{miralda2002} set a constraint of $\\sigma/m \\lesssim 0.02\\hbox{ cm}^2/\\hbox{g}$ on the velocity-independent elastic scattering cross section.  This constraint is one to two orders of magnitude tighter than other typical constraints on velocity-independent scattering \\citep{yoshida:weak:2000,gnedin2001,randall2008}.  It rendered velocity-independent scattering far too small to form cores in low surface brightness galaxies and other small galaxies \\citep{kuzio2008,deblok2008}. This is unfortunate because the main reason SIDM was interesting at the time was that it was a mechanism to create cores in such galaxies \\citep{spergel2000}.   The tightness of the SIDM constraints on cluster scales has meant that the focus of SIDM studies has shifted to those on velocity-dependent cross sections, such that SIDM may significantly alter dwarf-scale or smaller dark-matter halos while leaving cluster-mass halos largely untouched \\citep{feng2009,buckley2010b,loeb2011,vogelsberger2012}.  In recent times, such velocity-dependent interactions have arisen in hidden-sector models designed to interpret some charged-particle cosmic-ray observations  as evidence for dark-matter annihilation \\citep{pospelov2008,fox2009,arkanihamed2009,feng2010d}.  Constraints on other hidden-sector dark-matter models have been made using X-ray isophotes of the gas in the halo of the elliptical galaxy NGC 720 \\citep{buote2002,feng2009,feng2010d,buckley2010b,ibe2010,mcdermott2011,feng2012}.  However, as we show below, reports of the death of isotropic, velocity-independent elastic SIDM are greatly exaggerated. In this paper, we show that these earlier studies did not correctly account for the fact that the observed ellipticity (of the mass in cylinders or the projected gravitational potential) gets contributions from mass well outside the core, and the region outside the core retains its triaxiality. We also show that for ellipticity estimators that are relevant observationally, there is significant amount of scatter and the overlap between CDM and SIDM ellipticities is substantial even for $\\sigma/m=1\\hbox{ cm}^2/\\hbox{g}$. Lastly, we find that in the regions where SIDM particles have suffered (on average) about one or more interactions, the residual triaxiality is larger than what has been previously estimated \\citep{dave2001}.  Along with the analysis in Rocha et al. (2012), we find that studies of the central densities of dark-matter halos are likely to yield tighter constraints on the SIDM cross section than the morphology of the halos.  We briefly summarize our simulations in Sec. \\ref{sec:simulations}.  We present results on the three-dimensional shapes of  SIDM dark-matter halos compared to their CDM counterparts in Sec. \\ref{sec:3dshape}.  We reexamine the previous SIDM constraints based on halo shapes in light of our simulations in Sec. \\ref{sec:obs}.  In particular, we reexamine the \\citet{miralda2002} constraint in Sec. \\ref{sec:miralda} and from the shapes of the X-ray isophotes of NGC 720 \\citep{buote2002} in Sec \\ref{sec:xray}.  In Sec. \\ref{sec:lensing}, we show how other lensing data sets may constrain SIDM in the future.  We summarize the key points of this paper and present a few final thoughts in conclusion in Sec. \\ref{sec:conclusion}.   \\begin{table} \\label{sims.tab} \\centering \\noindent {\\bf Table 1:}  Summary of simulations.\\\\ \\begin{tabular}{|l|cccc|} \\hline Name & $L_\\mathrm{Box}$&  $m_\\mathrm{p}$ & $\\epsilon$ & $\\sigma/m$\\\\ &   [Mpc/h] &   [M$_\\odot$/h] & [kpc/h] &  [$\\hbox{cm}^2/\\hbox{g}$]  \\\\ \\hline \\hline CDM & $50$ &  $6.9\\times10^7$ & $1.0$ &  -- \\\\ & $25$ &  $8.6\\times10^6$ & $0.4$ &  -- \\\\ \\hline SIDM$_{0.1}$ & $50$  & $6.9\\times10^7$ & $1.0$ & 0.1\\\\ & $25$  & $8.6\\times 10^6$ & $0.4$  & 0.1\\\\ \\hline SIDM$_1$   & $50$ & $6.9\\times10^7$ & $1.0$ &  1.0\\\\ & $25$ &  $8.6\\times 10^6$ & $0.4$ &  1.0\\\\ \\hline \\end{tabular} \\vskip 0.5 cm *Note: columns give name, simulation box size, particle mass, force resolution, and interaction cross section. We use $h=0.71$.  See Rocha et al. (2012) for more details on the simulations. \\end{table}  \\begin{figure*} \\begin{center} \\includegraphics[width=1.0\\textwidth]{paper_SurfaceDensityMajor4_lowres.pdf} \\caption{\\label{fig:major4}Surface density of a halo of mass $M_{\\mathrm{vir}} = 1.2\\times 10^{14}M_\\odot$ projected along the major axis of the moment-of-inertia tensor -- the orientation that dominates the lensing probability.  The left column shows the halo for CDM, while the middle and right columns show the same halo simulated using SIDM with $\\sigma/m = 0.1$ $\\hbox{cm}^2/\\hbox{g}$ and $1.0$ $\\hbox{cm}^2/\\hbox{g}$, respectively.  The bottom row shows the same information, now zoomed in on the central region. The surface density stretches logarithmically from $\\approx 10^{-3}\\hbox{g}/\\hbox{cm}^2$ (blue) to $\\approx 10\\hbox{ g}/\\hbox{cm}^2$ (red).}% \\end{center} \\end{figure*}   \\begin{figure*} \\begin{center} \\includegraphics[width=0.95\\textwidth]{paper_SurfaceDensityIntermediate4_lowres.pdf} \\caption{\\label{fig:intermed4} Surface density profiles for the same halo shown in Fig. \\ref{fig:major4}, now projected along the intermediate axis.  Deviations from axisymmetry are highest along this projection.} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion} The takeaway message of this work is that mapping observations to constraints on the self-interaction cross section of dark matter is significantly more subtle than previously assumed, and as such, constraints based on halo shapes are, at present, one to two orders of magnitude weaker than previously claimed.  There are three primary reasons that contribute to this conclusion.  First, the observational probes (gravitational lensing and X-ray surface brightness) of halo shapes are actually probes of some moment of the mass distribution.  For lensing, the observational probes are also sensitive to all material along the line of sight.  While SIDM makes the three-dimensional density distribution significantly rounder within some inner radius $r$, the surface density will in general not be spherically symmetric at a projected radius $R=r$.  The surface densities are affected by material well outside the core set by scatterings where material is still quite triaxial (Fig. \\ref{fig:3drrvir}). Previous constraints were made under the assumption that the observations tracked the three-dimensional halo shape for fixed projected radius.  This is a less troublesome assumption for X-ray isophotes, since it is weighted by the square of the gas distribution, and hence sensitive to the central regions. The shapes measured should be related most closely to the shapes of the enclosed mass profile. So, to probe cross-sections as small as $\\sigma/m = 0.1  \\hbox{ cm}^2/\\hbox{g}$, one needs to get down to ${\\cal O}(10\\, \\rm kpc)$ from the center of the halos. The contribution of stars in this region makes it difficult to robustly estimate the shape of the dark matter profile and it also makes it difficult to get a large ensemble of galaxies for this study.  Second, there is a fair bit of scatter added by assembly history to the observed shapes and the scatter is large enough that it precludes using a small number of objects to set constraints on SIDM cross sections. Finally, although we find that the three-dimensional shape of halos begins to become more spherical than CDM at radii where the local interaction rate is fairly low, $\\Gamma(r) \\approx 0.1 \\, H_0$, there is a fair amount of triaxiality even when $\\Gamma(r) \\approx H_0$, a fact that was not appreciated in earlier studies \\citep[e.g.,][]{miralda2002,feng2010d}  We find that the convergence map of MS 2137-23 ($M_{\\mathrm{vir}} \\sim 10^{15}M_\\odot$) allows a velocity-independent SIDM cross section of $\\sigma/m = 1\\hbox{ cm}^2/\\hbox{g}$. The X-ray isophotes of NGC 720 ($M_{\\mathrm{vir}} \\sim 10^{13} M_\\odot$) likely rule out $\\sigma/m = 1\\hbox{ cm}^2/\\hbox{g}$, but are consistent with $\\sigma/m = 0.1\\hbox{ cm}^2/\\hbox{g}$ at radii where we can resolve shapes in our simulations.  Based on a preliminary comparison to lensing models of LoCuSS clusters, we conclude that $\\sigma/m = 0.1\\hbox{ cm}^2/\\hbox{g}$ is as consistent with observations as CDM but that $\\sigma/m = 1\\hbox{ cm}^2/\\hbox{g}$ is likely too large to be consistent with the observed shapes of those clusters.  Cross sections in this range are very interesting.    In Rocha et al. (2012), we show that a cross section in the neighborhood of $\\sigma/m = 0.1 \\hbox{ cm}^2/\\hbox{g}$ could solve the ``Too Big to Fail'' problem for the Milky Way dwarf spheroidals \\citep{boylan-kolchin2011b}, the core-cusp problem in LSB galaxies \\citep{kuzio2010,deblok2010}, as well as the shallow density profiles of the galaxy clusters in \\citet{sand2008,newman2009}, and \\citet{newman2011} while not undershooting their central densities or overshooting the core sizes.  Cross sections in this range are also consistent with other density-profile-based and subhalo-based constraints \\citep{yoshida:weak:2000,gnedin2001}.  Since the current set of observations appear to be consistent with a SIDM cross section of $\\sigma/m = 0.1\\hbox{ cm}^2/\\hbox{g}$, there are two relevant questions for shape-based SIDM constraints.  Will shape-based constraints be competitive with other types of SIDM constraints?  And what will it take to get down to $\\sigma/m \\sim 0.1\\hbox{ cm}^2/\\hbox{g}$ with shapes?  Upon closer inspection, our view is that constraints using existing data could be pushed below $\\sigma/m = 1\\hbox{ cm}^2/\\hbox{g}$, but it is not yet clear that we can get to $\\sigma/m \\sim 0.1\\hbox{ cm}^2/\\hbox{g}$.  While X-ray isophotes of the elliptical galaxy NGC 720 are consistent with $\\sigma/m = 0.1\\hbox{ cm}^2/\\hbox{g}$, there are some differences and a larger ensemble of elliptical galaxies may be able to test that.  There are a number of other elliptical galaxies for which high-resolution X-ray data exist \\citep[e.g.,][]{humphrey2006}, but they lack the detailed shape measurements of NGC 720. So better constraints could result from X-ray shape analysis for these galaxies.  For clusters, based on our quick pass through the LoCuSS results, lensing-based shape constraints on SIDM could also extend well below $\\sigma/m = 1\\hbox{ cm}^2/\\hbox{g}$ if simulations are performed of a statistically significant number of massive galaxy clusters. However, in studies of both galaxies and clusters, it is likely that the measured densities in the inner regions would be a better way to test for signatures of self-interacting dark matter."
    },
    "1208/1208.6329_arXiv.txt": {
        "abstract": "I present a study of high-resolution time series of \\CaIIH\\ images and \\FeI\\ $630.15~\\mathrm{nm}$ spectra taken with the Solar Optical Telescope on the Hinode spacecraft. There is excellent correspondence between the \\CaIIH\\ and \\FeI\\ line core intensity, except tenuous emission around the network field concentrations in the former that is absent in the latter. Analysis of on-disk observations and a comparison with limb observations suggests that this ``network haze'' corresponds to spicules, and likely to type-II spicules in particular. They are known to appear in emission in on-disk broadband \\CaIIH\\ diagnostics and the network haze is strongest in those areas where features similar to type-II spicules are produced in simulations. ",
        "introduction": "\\label{sec:introduction}  Spicules, mottles, and fibrils have long been studied as the principal component of the magnetized solar chromosphere. The review by \\cite{1968SoPh....3..367B} on spicules at the limb and more recent work on mottles and fibrils on the disk \\cite[e.g.,][]{2006ApJ...647L..73H,2007ApJ...655..624D} all emphasize the crucial role of these features in the mass and energy balance of the solar chromosphere. The relationship between spicules and similar features on disk has not been clearly established, largely because of the confusion resulting from projection as well as their fine structure and fast dynamics.  \\cite{2007PASJ...59S.655D} argue that at least two species of spicules exist. Both are tied intrinsically to magnetic field, but are different in their dynamics. The traditional ``type-I'' spicules exhibit slower evolution ($3$--$7~\\mathrm{min}$) and up-and-down motions with velocities on the order of $20~\\mathrm{km}\\,\\mathrm{s}^{-1}$. The ``type-II'' spicules, first identified as ``straws'' in near-limb \\CaIIH\\ images from the Dutch Open Telescope \\citep{2004A&A...413.1183R} by \\cite{2006ASPC..354..276R,2007ASPC..368...27R}, show much shorter lifetimes ($10$--$60~\\mathrm{s}$), faster velocities ($50$--$150~\\mathrm{km}\\,\\mathrm{s}^{-1}$), and often disappear over their whole length within a few seconds. Further studies of type-II spicules have revealed that they are associated with Alfv\\'enic waves \\citep{2007Sci...318.1574D,2011Natur.475..477M}.  While ``straws'' or type-II spicules are clearly evident at or near the limb, they are not so readily identified on the disk. Their faint, tenuous appearance makes them hard to discern against the background. A search for disk counterparts by \\cite{2008ApJ...679L.167L} and a more detailed study by \\cite{2009ApJ...705..272R} revealed the existence of ``rapid blueshift excursions'' (RBEs) in the \\Ha\\ and \\CaII\\ IR lines that exhibit similar dynamics to type-II spicules and fit well with the finding by \\cite{2009ApJ...701L...1D} and \\cite{2009ApJ...706L..80M} that type-II spicules are possibly connected to the transition region and the corona through the blue-wing emission excess found in EUV lines by \\cite{2008ApJ...678L..67H} \\citep[see also][]{2010ApJ...718.1070H, 2010ApJ...722.1013D, 2011Sci...331...55D, 2011ApJ...732...84M, 2011ApJ...727L..37T, 2011ApJ...738...18T, 2012ApJ...746..158J}.  The correspondence of RBEs and type-II spicules seems well-established \\citep{2012ApJ...752..108S}. However, type-II spicules have been exclusively identified from Hinode images, and RBEs have only been identified in ground-based data from narrowband imaging instruments, so a direct identification of the on-disk counterpart of type-II spicules has yet to be presented. In this Letter, I identify a weak haze around the magnetic network using Hinode data that I identify as a candidate for the on-disk manifestation of type-II spicules. ",
        "conclusions": "\\label{sec:discussion}  There are three contributions to \\CaIIH\\ network brightness in observations with relatively broad filter passbands such as these that cause an excess over the internetwork brightness. Because the filter is broad compared to the core of the \\CaIIH\\ line the majority of the intensity in these observations is photospheric sampled by the line wings. The classic Wilson depression allows emission to escape from deeper and hotter layers and hence the brightness of the wings is increased \\citep[e.g.,][]{2005A&A...437.1069S}. Calculations of the response function by \\cite{2007PASJ...59S.663C} show that the passband also samples higher layers. Heating in the magnetic network in the high photosphere and low chromosphere thus adds to the network brightness excess \\citep[e.g.,][]{2009A&A...503..577C}. Finally, there is a contribution from chromospheric structures in the core of the \\CaIIH\\ line \\citep[][]{2009A&A...500.1239R}.  The core of the \\FeI\\ line samples the same atmospheric layers as the Hinode \\CaIIH\\ passband as is evidenced by the exceedingly high correlation coefficient between the two diagnostics. As a result, the first two sources of brightness excess in the \\CaIIH\\ images will also contribute in the same way to the excess of brightness in the network as seen in the core of the \\FeI\\ line. The third source, however, is mostly absent in the \\FeI\\ observations because of the lack of optical thickness of spicules in that line. A weak emission signal extending less than $1\\arcsec$ above the limb was recently identified by \\cite{2010ApJ...713..450L}. They analyze the emission and conclude that the polarization signature is most likely created by scattering and that it is depolarized through the Hanle effect by a magnetic field of several Gauss.  Limb observations show that the distinct \\CaIIH\\ spicule forest is not present in the \\FeI\\ core intensity. This, and the correspondence of the length of the spicules in \\CaIIH\\ to the extent of the on-disk network haze (see Fig.~\\ref{fig:haze}), suggests that the haze is caused by spicules.  The persistence of some arch-like structures on the limb is noteworthy. The chromospheric magnetic field, stable on timescales longer than a few minutes, is apparently traced out by spicules preferentially along some field lines.  Type-I spicules have been extensively studied in \\CaIIH\\ \\citep{1968SoPh....3..367B}. Mottles and fibrils, on the other hand, have not been identified in \\CaIIH\\ diagnostics. They have been studied in chromospheric observations in lines such as \\Ha\\ and the IR \\CaII\\ triplet. It stands to reason that they must have some signature also in the \\CaIIH\\ line, but they probably appear as absorption features that would be difficult to detect against the background of reversed granulation.  It is more likely that the network haze is the on-disk counterpart of type-II spicules. Type-II spicules have been detected on disk as emission ``straws'' by \\cite{2006ASPC..354..276R,2007ASPC..368...27R}. They cannot be individually identified in these \\CaIIH\\ images (neither in the composite images shown in Figs.~\\ref{fig:data_20070419} and~\\ref{fig:data_20070219} nor in the individual images of which they are made up). This can be attributed to several factors. First, the fast dynamics of the very slender type-II spicules are not captured in these \\CaIIH\\ images because they are resampled in time to match the slow scanning of the SP slit. Furthermore, the wide SOT \\CaIIH\\ bandpass of $0.3~\\mathrm{nm}$ samples too much light in the wings of the line for the type-II spicules to show up with high contrast in disk observations. It is possible to observe type-II spicules directly on the limb with SOT because there is no bright background \\citep[e.g.,][]{2007PASJ...59S.655D}, and even near the limb \\citep{2010ApJ...714L...1S}, though in both cases there is confusion from integration along the line of sight. In SOT \\CaIIH\\ filtergrams, type-II spicules would be very faint on disk, yet must have some emission signature. RBEs, on the other hand, appear in absorption against the photosphere in \\Ha\\ because that line has no opacity in the region of the atmosphere where reversed granulation is formed \\citep{2006A&A...449.1209L}.  The \\CaIIH\\ intensity excess is strongest in the region of emerging flux round ($20\\arcsec$,$-15\\arcsec$). It is however also present and faintly visible around other areas of concentrated flux, e.g., the network vertex at ($0\\arcsec$,$5\\arcsec$). The average intensity as a function of distance to strong flux concentrations still shows excess \\CaIIH\\ emission between $1\\arcsec$ and $5\\arcsec$ even if the emerging flux region is excluded from the analysis. This is in agreement with the results of \\cite{2012ApJ...752..108S} who studied RBEs and found they are ubiquitous but concentrated around regions of enhanced magnetic field. From simple tests it appears that the results presented here are robust with respect to the choice of threshold value for the selection of concentrations of strong vertical flux. If the region of emerging flux is excluded, the \\FeI\\ intensity shows an excess over the \\CaIIH\\ intensity for distances less than $1\\arcsec$ from strong vertical flux, consistent with the observation in Fig.~\\ref{fig:data_20070419} that the flux concentrations themselves are bright in the difference image.   The simulations by \\cite{2011ApJ...736....9M} also provide evidence to support the identification of the haze with type-II spicules. They find structures in their realistic 3D MHD simulations that resemble type-II spicules in predominantly unipolar regions of emerging flux. The difference image in Fig.~\\ref{fig:data_20070419} shows that the network haze is strongest in the two patches of predominantly unipolar flux of opposite polarity in the region of enhanced activity around ($28\\arcsec$,$-18\\arcsec$) and ($33\\arcsec$,$-10\\arcsec$), in agreement with the simulations. The objects in the simulation are produced as a result of deflection of plasma that is pushed against a ``wall'' of low-$\\beta$ plasma by a strong Lorentz force. The regions of emerging flux exhibit the strong gradients in the magnetic field in both their simulation and these observations, which is required for the mechanism to work effectively. Ideally, for comparison with these simulations one would not only determine the locations of network haze but also the history of flows in these areas through feature tracking on a sequence of magnetograms. The observations used here unfortunately did not include such a sequence.  Because of the relatively broad SOT \\CaIIH\\ filter, the wings of the line dominate the contribution to the reversed granulation intensity. If a diagnostic can be found that samples similar structures as the \\CaIIH\\ wing but that lacks opacity in chromospheric structures, it may be possible to increase the contrast of spicules on disk through subtraction. Such a procedure has previously been used for the identification of magnetic bright points \\citep{1998ApJ...509..435V}. The \\FeI\\ line core would appear to be a suitable candidate. However, the filter bandwidth must be narrow. The SP samples the \\FeI\\ line very cleanly with a spectral resolution of $3.0~\\mathrm{pm}$ sampled with $2.15~\\mathrm{pm}$ pixels. The difference image in Fig.~\\ref{fig:data_20070219} suffers from the spatial resolution of the SP data as well as the resampling of the \\CaIIH\\ data in time to match slit scanning.  Summarizing, there is evidence to suggest that the network haze in the \\CaIIH\\ images is the on-disk manifestation of some class spicules. Type-II spicules are the likely candidate, since they are known to appear in emission in on-disk \\CaIIH\\ images from other other telescopes \\citep{2006ASPC..354..276R,2007ASPC..368...27R}, and the region around which is appears most clearly is where type-II spicules are expected to be produced based on simulations \\citep{2011ApJ...736....9M}."
    },
    "1208/1208.4606_arXiv.txt": {
        "abstract": "{ We present a self-consistent three-dimensional Monte-Carlo radiative transfer model of the stellar and dust emission in the Milky-Way, and have computed synthetic observations of the 3.6 to 100\\microns emission in the Galactic mid-plane. In order to compare the model to observations, we use the GLIMPSE, MIPSGAL, and IRAS surveys to construct total emission spectra, as well as longitude and latitude profiles for the emission. The distribution of stars and dust is taken from the SKY model, and the dust emissivities includes an approximation of the emission from polycyclic aromatic hydrocarbons in addition to thermal emission. The model emission is in broad agreement with the observations, but a few modifications are needed to obtain a good fit. Firstly, by adjusting the model to include two major and two minor spiral arms rather than four equal spiral arms, the fit to the  longitude profiles for $|\\ell|>30^\\circ$ can be improved. Secondly, introducing a deficit in the dust distribution in the inner Galaxy results in a better fit to the shape of the IRAS longitude profiles at 60 and 100\\microns. With these modifications, the model fits  the observed profiles well, although it systematically under-estimates the 5.8 and 8.0\\microns fluxes. One way to resolve this discrepancy is to increase the abundance of PAH molecules by 50\\% compared to the original model, although we note that changes to the dust distribution or radiation field may provide alternative solutions. Finally, we use the model to quantify which stellar populations contribute the most to the heating of different dust types, and which stellar populations and dust types contribute the most to the emission at different wavelengths. } ",
        "introduction": "\\label{sec:introduction}  Although it has been common to characterize the large-scale Galactic diffuse emission by treating each component of the Galaxy individually -- the neutral gas distribution, HII region maps, the magnetic field structure -- this emission is usually caused by a convolution of different combinations of these components. Far-infrared dust emission, for example, traces the dust distribution heated by stars; radio free-free emission depends specifically on the local density of main-sequence massive stars; radio synchrotron emission traces the local cosmic ray electron density (and energy spectrum) convolved with the magnetic field structure; and so on. Since our knowledge of the distribution of any one of these components is still rather incomplete, modeling emission that arises from the coupling of two or more of them is a daunting task, requiring numerous assumptions to develop a complete model that can be compared with observations.  The value of such an exercise lies in identifying the factors that have the largest influence on determining the properties of the observed emission, with the hope that eventually {\\it all} of the different diffuse emission measurements can be combined to yield a more complete picture of the Galaxy.  In this paper, we present models of the stellar and diffuse infrared emission in the Galactic plane. Although investigations of the diffuse infrared emission have been carried out in the past, to our knowledge this is the first work to focus on the mid-infrared (3.6-8.0 $\\mu$m) diffuse emission from polycyclic aromatic hydrocarbons (PAHs), which is much more sensitive to the distribution of the young stars than the longer wavelength dust emission probed by IRAS and COBE/DIRBE.  The large-scale distribution of stars and diffuse dust emission in the Milky-Way has been the subject of numerous studies. On one hand, powerful three-dimensional stellar distribution models have been developed to simulate the inventory of stellar populations as a function of location in the Galaxy and position on the sky. The Besan\\c{c}on model \\citep{robin:86:71, bienayme:87:94, robin:96:125, robin:03:523},  the SKY model \\citep{wainscoat:92:111, cohen:93:1860, cohen:94:582, cohen:95:874}, and the TRILEGAL model \\citep{girardi:05:895} are three of the most widely known such models aimed at reproducing the stellar populations of the Galaxy. Such models typically include various components to describe the Galaxy, such as a bulge, halo, disk(s), spiral arms, and so on, and include stars at various masses, metallicities, and stages of evolution.  On the other hand, analytic models of the diffuse dust emission in the Galaxy have also been developed to reproduce the large-scale mid- and far-infrared emission. For example, \\citet{sodroski:97:173} used COBE/DIRBE data from 12$-$240\\microns to determine the best-fitting abundances and temperature of large dust grains, the abundance of very small transiently heated dust grains and PAHs, and the energy density of the interstellar radiation field as a function of Galactocentric radius. Similarly, \\citet{davies:97:679} fit a model for the emission from cool (18-22K) dust to the COBE/DIRBE far-infrared data. More recently, the \\citet{planck:11:A21} fit the observations of the Galaxy by the Planck satellite, using a model similar to that used by \\citeauthor{sodroski:97:173}  \\citet{drimmel:00:L13} carried out a qualitative comparison of COBE/DIRBE K-band and 240\\microns observations, and found that while the near-infrared shows evidence only for two stellar arms, the far-infrared observations are consistent with four arms (in agreement with the influential HII region study of \\citealt{georgelin:70:349}), leading to the conclusion that the main Galactic potential is two-armed, but that the structure of the dust and gas responding to the potential is more complex, and can be adequately characterized as four-armed. \\citet{drimmel:01:181} extended this by developing a quantitative model to fit the above data, which consisted of a parametric model of the dust distribution and associated far-infrared emission. This dust distribution was then used by \\citet{drimmel:03:205} to predict the three-dimensional K-band extinction.  What has been lacking in these models however is a fully self-consistent treatment of the {\\it source} of the dust heating -- the photons produced by stars of various temperatures and luminosities -- and the {\\it target} of the heating -- grains of different sizes and PAHs -- distributed throughout the disk of the Galaxy. This was addressed in particular by \\citet{porter:05}, who developed a self-consistent radiative transfer model of the stellar and dust emission in order to study the propagation of cosmic-rays in the Galaxy, and computed the interstellar radiation field as a function of position in the Galaxy, assuming azimuthal symmetry. The model was subsequently used and refined in \\citet{moskalenko:06:L155}, \\citet{porter:06:L29}, and \\citet{porter:08:400}, and was successfully compared to the local integrated all-sky interstellar radiation field as seen from the Sun.   In this paper we present a fully self-consistent 3-dimensional radiative transfer model of the Galaxy at infrared wavelengths, considering both stellar and dust emission, and we compare the model to the distribution of emission on the sky from near-infrared to far-infrared wavelengths. The model is self-consistent in the sense that the heating of the dust does not follow an analytical prescription, but is instead directly computed from the stellar populations, similarly to the \\citeauthor{porter:05} model. The main aims of this paper with this model are the following:   \\begin{itemize} \\item To determine whether an existing Galactic model for stellar populations and dust -- here the SKY model -- used in conjunction with the dust properties determined by \\citet{draine:07:810}, which include transiently heated very small grains and PAH molecules, adequately reproduces the stellar and diffuse emission seen at mid- and far-infrared wavelengths. \\item To determine whether any modifications are needed to this model to reproduce the observations, and whether any new insights on Galactic structure can be gained. However, we stress that the aim is not to provide an exhaustive parameter space study, as given the number of stellar populations and components, the number of parameters describing the model is of the order of several hundred. Rather, the aim is to determine possible and realistic modifications that improve the model. \\item To determine the relative contributions of resolved and unresolved stellar flux, emission from PAH molecules and dust grains, and scattering at infrared wavelengths. \\item To determine the relative importance of various stellar populations in heating dust grains and exciting PAH molecules. \\end{itemize}  In Section \\S\\ref{sec:observations} we describe the observational datasets used in this paper. In \\S\\ref{sec:model}, we describe the initial model used, the dust properties, and the radiative transfer code. In \\S\\ref{sec:results}, we show results for both the initial model (\\S\\ref{sec:initial}) as well as an improved model (\\S\\ref{sec:improved}). In \\S\\ref{sec:analysis} we examine the main contributors to the observed flux and to the heating of the dust (\\S\\ref{sec:contributions}), we study how much IRAC stellar flux is likely to be unresolved (\\S\\ref{sec:unresolved}), and we generate images of the Galaxy from an external viewpoint (\\S\\ref{sec:external}). Finally, in \\S\\ref{sec:summary} we summarize our findings. ",
        "conclusions": "\\label{sec:summary}  A radiative transfer model of the Galaxy which uses the SKY model in conjunction with  the dust properties from \\citet{draine:07:810} was developed to self-consistently calculate the heating of dust grains, and verify whether it is able to reproduce the observed surface brightness from 3.6\\microns to 100\\microns. The main findings presented in this paper are the following:  \\begin{enumerate} \\item The initial model is able to roughly reproduce the order of magnitude of the observed surface brightnesses observed, but there are disagreements between the model and observations, notably for $|\\ell| > 30^\\circ$ for the IRAC and MIPS bands, and inside $|\\ell| < 30^\\circ$ for IRAS 60\\microns and 100\\microns. \\item By modifying the model to incorporate two major stellar spiral arms and two secondary spiral arms with only young massive stars, as well as removing dust from the central few kpc of the Galaxy, we are able to significantly improve the quality of the fit. A slight systematic offset remains at IRAC 5.8 and 8.0\\microns. By increasing the abundance of PAHs by 50\\%, we are able to eliminate this systematic offset, though we caution that other effects may explain this offset. \\item The flatness of the IRAC 5.8 and 8.0\\microns and MIPS 24\\microns emission is not directly due to a hole in the dust distribution, but is a consequence of the lack of strong UV sources within the inner few kpc of the Galactic center. The evidence for a deficit of dust is only apparent at longer wavelengths, in the IRAS 60 and 100\\microns bands. \\item Since the stellar populations are not symmetrically distributed, the heating of the dust is not uniform, and therefore the spiral arms appear in diffuse emission at infrared wavelengths without the need for a non-symmetrical dust distribution (i.e. the model does not require additional dust in the spiral arms). This is not to say that the distribution of dust is not enhanced in the spiral arms, but that there is no evidence from the data that it is. \\item The overall flux at IRAC 3.6 and 4.5\\microns is dominated by giant stars, PAHs, and main-sequence stars, while the flux from IRAC 5.8\\microns to IRAS 100\\microns probes mostly dust, with the size of the dust grains probed increasing with wavelength. \\item On large scales, transiently heated very small grains and PAH molecules are predominantly heated by B-type stars. The larger dust grains also have a significant component of heating from giant stars. \\item The \\textit{Spitzer}/IRAC bands may contain as much unresolved stellar flux as resolved stellar flux for $|\\ell| < 20^\\circ$. \\end{enumerate}  The models presented here, while simple, help us quantify the main factors that determine the large-scale Galactic mid-plane emission. We plan to carry out significant improvement to the models in the future, including a more systematic exploration of parameter space, a more realistic clumpy distribution of dust, and an improved treatment of the inner Galaxy, specifically relating to the molecular ring and the Galactic bar. In addition, the modeling of observations can be extended to spatial and wavelength regions outside that covered here. For instance, one could include observations for the outer Galaxy (GLIMPSE 360) or at longer wavelengths (\\textit{Herschel} HiGal). Once the data from the Wide-Field Infrared Survey Explorer (WISE) are fully released, it will even be possible to model the all-sky observations from 3.5\\microns to 100\\microns when combined with the IRAS all-sky data."
    },
    "1208/1208.5436_arXiv.txt": {
        "abstract": "{The Atacama Large Millimeter/submillimeter Array (ALMA) will have the necessary resolution to observe a planetary gap created by a Jupiter-mass planet in a protoplanetary disk. Because it will observe at submillimeter and millimeter wavelengths, grains in the size range 10\\,$\\mu$m to 1\\,cm are relevant for the thermal emission. For the standard parameters of a T~Tauri disk, most grains of this size range are weakly coupled to the gas (leading to vertical settling and radial migration) and the common approximation of well-mixed gas and dust does not hold.} {We provide predictions for ALMA observations of planet gaps that account for the specific spatial distribution of dust that results from consistent gas+dust dynamics.} {In a previous work, we ran full 3D, two-fluid Smoothed Particle Hydrodynamics (SPH) simulations of a planet embedded in a gas+dust T~Tauri disk for different planet masses and grain sizes. In this work, the resulting dust distributions are passed to the Monte Carlo radiative transfer code \\mcfost\\ to construct synthetic images in the ALMA wavebands. We then use the ALMA simulator to produce images that include thermal and phase noise for a range of angular resolutions, wavelengths, and integration times, as well as for different inclinations, declinations and distances. We also produce images which assume that gas and dust are well mixed with a gas-to-dust ratio of 100 to compare with previous ALMA predictions, all made under this hypothesis.} {Our findings clearly demonstrate the importance of correctly incorporating the dust dynamics. We show that the gap carved by a 1\\,$M_\\mathrm{J}$ planet orbiting at 40\\,AU is visible with a much higher contrast than the well-mixed assumption would predict. In the case of a 5\\,$M_\\mathrm{J}$ planet, we clearly see a deficit in dust emission in the inner disk, and point out the risk of interpreting the resulting image as that of a transition disk with an inner hole if observed in unfavorable conditions. Planet signatures are fainter in more distant disks but declination or inclination to the line-of-sight have little effect on ALMA's ability to resolve the gaps.} {ALMA has the potential to see signposts of planets in disks of nearby star-forming regions. We present optimized observing parameters to detect them in the case of 1 and 5\\,$M_\\mathrm{J}$ planets on 40\\,AU orbits.} ",
        "introduction": "\\label{sec:Introduction}  The study of exoplanets has steadily gained momentum in the past two decades, in particular since the confirmed identification of the first exoplanet around the solar-like star 51~Peg \\citep{MayorQueloz95}. More than 700 planets located in $\\sim$500 planetary systems are now confirmed\\footnote{\\texttt{http://exoplanet.eu}}. A direct consequence of these surveys is that more reliable estimations of the statistical properties of planetary systems are becoming available (e.g., distributions of orbital distances and planet masses, see \\citealp{Schneider2011}).  Two scenarios currently compete for giant gaseous planet formation: core-accretion \\citep[e.g.][]{AMBW05} and disk gravitational instability \\citep[e.g.][and references therein]{Boss2011}, each with its own merits and limitations. In the core-accretion scenario, giant planets form in two steps. A rocky core is first formed by progressive accumulation of planetesimals until a critical mass ($\\sim$10\\,$M_\\oplus$) is reached. Then rapid capture of nebular gas can occur, in a runaway fashion, until the planet reaches its final mass. In the disk instability scenario, a gravitationally unstable disk fragments into self-gravitating clumps which then directly form gas giant planets. The relevance of each mechanism to explain the current observations of exoplanets, especially at large orbital radii, depends largely on the exact properties of the underlying protoplanetary disk \\citep[see e.g.][]{Mordasini2012}.  Although circumstellar disks of gas and dust were imaged directly roughly at the same time as the exoplanets \\citep{Dutrey1994, Burrows1996} it is only recently that the link between them and the exoplanets was confirmed directly by observations when planets and disks were imaged simultaneously around $\\beta$~Pictoris \\citep{Lagrange2010} and HR~8799 \\citep{Marois2008,Marois2010}. Both systems are however rather evolved and gas giant formation has likely ceased in these now gas-poor debris disks. Interestingly, the case of HR~8799 with several planets orbiting between 14 and 68\\,AU suggests that the formation of giant planets on wide orbits is a strong possibility. Indeed, \\citet{Quanz2012} recently found that a larger fraction of stars than previously assumed may harbor gas giant planets at larger orbital separations and \\citet{Lambrechts2012} showed that such planets can be formed by the core accretion mechanism well within the disk lifetime.  ALMA, when completed, will have the capacity to observe these wide planetary systems at a younger age and explore a parameter space that is complementary to that probed by radial velocity and transit techniques. In this paper we show that signs of the presence of these planets on wide orbits, and specifically the gaps they carve in the disk, will be detectable with ALMA. The process of gap formation depends on the planet mass, the disk surface density and the size of dust grains \\citep{PM04,Fouchet07} and measuring their width and depth can provide constraints on the underlying disk structure \\citep{CMM06}. Previous predictions indicate that ALMA should be able to detect such gaps \\citep{W02,W05}. Most of these predictions were made assuming the very best performance ALMA will provide, i.e., longest baselines, shortest wavelengths, no or low phase noise.  The dust responsible for the emission was also assumed to be well-mixed with the gas, irrespective of the dust particle size. Here we will use the results of numerical simulations \\citep[hereafter Paper~I]{Fouchet10} where the dust distributions are calculated for each grain size using a two-phase SPH approach to produce ALMA synthetic images. We will assess the different planet signatures due to different dust distributions as a function of the grain size. We will explore the ability of ALMA to detect these gaps, but also to measure the disk properties both inside and outside the planet's orbit for various combinations of observing time, angular resolution and wavelength in different star forming regions in order to estimate the optimum survey strategy.  \\defcitealias{Fouchet10}{Paper~I} ",
        "conclusions": "\\label{sec:Conclusion}  In a previous study \\citepalias[see][]{Fouchet10}, we have run 3D hydrodynamic simulations of the gaps carved in a 0.02\\,$M_\\odot$ disk of gas and dust by 1 and 5\\,$M_\\mathrm{J}$ planets at 40\\,AU from the 1\\,$M_\\odot$ central star. We followed consistently the dynamics of dust grains in the range of sizes contributing the most to the ALMA wavelengths and described the distinctive, sharper features in the dust phase. In this paper, we now provide predictions of observations with ALMA and examine the detectability of these planetary gaps. We first produce raw synthetic images by applying the radiative transfer code \\mcfost\\ to the hydrodynamic results. We produce two kinds of images: the ``dynamic'' case where we use the true spatial distribution of dust of different grain sizes obtained from the hydrodynamic simulations and the ``well-mixed'' case where we simply use the spatial distribution of gas and assume that the dust follows it in order to check the validity of the ``well-mixed'' hypothesis.  We then produce simulated ALMA images using the CASA software package. We investigate the choice of integration time, wavelength and angular resolution. We show that an integration time of 1\\,h is enough to ensure a firm detection of our planet gaps and that the gain in signal-to-noise as we increase the integration time does not warrant the necessary extra observing time on a facility that will be heavily oversubscribed. We also show that an angular resolution of $0.10''$ is optimal to resolve the gap with the best contrast in most nearby star-forming regions. We then show that the use of small or large wavelengths combined with these optimal observing parameters give poor to bad results. In average sky conditions, we find the wavelength giving the best signal-to-noise ratio to be 850\\,$\\mu$m when phase noise is neglected or 1.3\\,mm, with slightly lower performance, when it is included. However, the use of WVRs should provide real-time correction of phase noise and favor 850\\,$\\mu$m as the wavelength of choice. In the driest weather, phase noise is negligible.  A number of previous studies of planet gaps \\citep[e.g.][]{W02,W05}, and even recent work on Rossby vortices in protoplanetary disks \\citep{Regaly2012}, have relied on hydrodynamic simulations of 2D gas-only disks to produce simulations of ALMA observations by reconstructing a three-dimensional distribution of dust that follows that of the gas in the midplane and is in hydrostatic equilibrium in the vertical direction before feeding it to a radiative transfer code. In addition to neglecting the vertical dust settling, this approach also results in smoother dust features, leading to more pessimistic predictions for the detectability of the disk features. The comparison of our simulated images in the ``well-mixed'' and ``dynamic'' cases show that they are very different and that the more realistic case of including a self-consistent treatment of the dust dynamics produces much sharper features and more clearly defined gaps. We would like to stress here that investigators need to be cautious when assessing the observability of a particular structure in their source from approximate methods and that a procedure similar to our ``dynamic'' case should be used whenever possible.  With our ``dynamic'' approach, we find that the gap carved by a 1\\,$M_\\mathrm{J}$ planet orbiting 40\\,AU from its star is easily detected. Images for the 5\\,$M_\\mathrm{J}$ planet show a high contrast between the outer disk and the deeper and wider gap, but the fainter inner disk is barely seen with our standard observing parameters. In order to detect the regions interior to the gap and remove the possible confusion with transition disks with large inner holes, one needs to observe at shorter wavelengths in the driest conditions possible or use a combination of longer integration time and higher angular resolution. In any case, multi-wavelength followup observations of detected gaps is advisable for a characterization of the planet.  This work demonstrates that the gap carved by a moderately massive planet of one Jupiter mass at large orbital radii will be well within reach of ALMA. While the image quality naturally decreases with increasing distance, we find that the ability to observe planet gaps is little affected by the source declination or inclination to the line-of-sight. Although few planets on wide orbits have been detected so far due to the biases of current techniques, their existence is supported by the more recent planet formation models as well as by observational constraints on planet populations. ALMA, probing a new part of the orbital parameter space, has the potential to unveil more of them in nearby star-forming regions."
    },
    "1208/1208.2350_arXiv.txt": {
        "abstract": "We present a map of the Cosmic Microwave Background (CMB) anisotropies induced by the late Integrated Sachs Wolfe effect. The map is constructed by combining the information of the WMAP 7-yr CMB data and the NRAO VLA Sky Survey (NVSS) through a linear filter. This combination improves the quality of the map that would be obtained using information only from the Large Scale Structure data.  In order to apply the filter, a given cosmological model needs to be assumed. In particular, we consider the standard $\\Lambda$CDM model. As a test of consistency, we show that the reconstructed map is in agreement with the assumed model, which is also favoured against a scenario where no correlation between the CMB and NVSS catalogue is considered. ",
        "introduction": "\\label{intro}  Recent observations, as cosmic microwave background (CMB), supernovae type Ia (SNIa) or baryon acoustic oscillations (BAOs) agree in establishing a current accelerated expansion of the Universe~\\citep[see][for a recent review]{weinberg2012}, which, despite some less popular interpretations, is believed to be caused by the presence of some dark energy~\\citep[see][for a review]{peebles2003}. The nature of dark energy is one of the most puzzling issues in modern cosmology. The actual characteristics of this fluid are still unclear, although, up to date, a good agreement is found between the observations and the predictions derived from the presence of a cosmological constant with an equation of state $p=-\\rho$.  One of the classical probes of dark energy is given by a non-null contribution to the CMB anisotropies from the late integrated Sachs-Wolf effect~\\citep[ISW,][]{sachs1967}, under the assumption of a spatially flat universe. The (linear) ISW fluctuations are higher at very large angular scales, but, in any case, much smaller than the primary CMB fluctuations. In a seminal work,~\\cite{crittenden1996} proposed the cross-correlation of the CMB fluctuations and the dark matter distribution (typically traced by galaxy catalogues) as a possible approach to detect the ISW. Soon after the release of the WMAP data,~\\cite{boughn2004} reported the first detection of the ISW effect via the CMB and the galaxy number density field. Posterior works \\citep[e.g.,][]{fosalba2003,vielva2006,pietrobon2006,ho2008,giannantonio2008,mcewen2008,dupe2011,schiavon2012} have confirmed the detection of the ISW effect by exploring several galaxy catalogues and cross-correlation techniques. Average detection is found at $\\approx 3\\sigma$.  Besides these works focused on the statistical detection of the ISW, more recently there have been attempts to recover the actual ISW fluctuations on the sky.  In principle, an optimal ISW map could be derived from a 3D gravitational potential. This has been explored from very large simulations~\\citep[e.g.,][]{cai2010}. The complexity to recover optimally the potential from surveys of galaxies with redshift information is challenging~\\cite[e.g.,][]{kitaura2009,jasche2010}, but very promising from the ISW studies point of view~\\cite[e.g.,][]{frommert2008}.  Other approaches can be followed from surveys where the redshift information is poor, or known only statistically.~\\cite{barreiro2008} proposed to use jointly maps of CMB anisotropies and of the galaxy number density field to recover the ISW signal on the sky. Alternative works making use only of galaxy catalogue maps have been proposed afterwards: \\cite{granett2009} on LRGs from SDSS-DR6, or~\\cite{francis2010a} and~\\cite{dupe2011} on 2MASS.  This paper presents an application of the approach described in~\\cite{barreiro2008} to WMAP \\citep{jarosik2011} and NVSS~\\citep{condon1998}. The outline of the article is as follows. The methodology is reviewed in Section~\\ref{sec:method}. A description of the data used, as well as the fiducial theoretical model is presented in Section~\\ref{sec:data}. In Section~\\ref{sec:analysis} we present the results. Finally, conclusions are given in Section~\\ref{sec:final}. ",
        "conclusions": "\\label{sec:final} We have presented a reconstructed map of the ISW effect, obtained by combining through a linear filter CMB and LSS data. In particular, we have used the 7-yr WMAP data and the NVSS galaxies catalogue. The joint combination of both data sets improves by a 15 per cent the error of the ISW reconstruction in comparison to the case when only the NVSS data are used. We have performed a consistency test, showing a good agreement between the cross correlation inferred from the reconstructed ISW map and the assumed fiducial model. In particular, the data favour a $\\Lambda$CDM model with respect to a scenario with null correlation between the CMB and the NVSS data. The relative contribution to the angular power spectrum of the ISW reconstructed map is dominated by the CMB fluctuations up to $\\ell \\approx 10$ and by the density number galaxies field at larger multipoles.   The presented methodology works in harmonic space, which implies the use of surveys with large sky coverage, in order to avoid the problematics introduced by large masks. However, this technique can be easily extended to work directly in the pixel space, which would make straightforward to deal with a mask and, in particular, LSS surveys with smaller sky coverage could also be used to reconstruct the ISW signal~\\citep[][in preparation]{bonavera2012}. In addition, several catalogues can be combined at the same time, provided the covariance matrix between the surveys and the CMB data is known.  Finally, let us remark that the application of the approach described in this paper to future surveys (as EUCLID~\\citealt{refregier2010} or J-PAS~\\citealt{benitez2009}, with very large sky coverage and very accurate redshift estimation) could provide maps of the ISW anisotropies caused by the large-scale structure at different redshift shells. This will provide a tomographic view of the ISW fluctuations."
    },
    "1208/1208.2485_arXiv.txt": {
        "abstract": "Recently, much attention has been given to double-peaked narrow emission-line galaxies, some of which are suggested to be related with merging galaxies.  We make a systematic search to build the largest sample of these sources from Data Release 7 of Sloan Digital Sky Survey (SDSS DR7). With reasonable criteria of fluxes, full-width-half-maximum of emission lines and separations of the peaks, we select 3,030 double-peaked narrow emission-lines galaxies. In light of the existence of broad Balmer lines and the locations of the two components of double-peaked narrow emission lines distinguished by the Kauffmann et al. (2003) criteria in the Baldwin-Phillips-Terlevich (BPT) diagram, we find that there are 81 type I AGN, 837 double-type II AGN (2-type II), 708 galaxies with double star forming components (2-SF), 400 with mixed star forming and type II AGN components (type II + SF) and 1,004 unknown-type objects. As a by-product, a sample of galaxies (12,582) with asymmetric or top-flat profiles of emission lines is established. After inspecting the SDSS images of the two samples visually, we find 54 galaxies with dual cores. The present samples can be used to study the dynamics of merging galaxies, the triggering mechanism of black hole activity, the hierarchical growth of galaxies and dynamics of narrow line regions driven by outflows and rotating disk. ",
        "introduction": "Double-peaked narrow emission-line active galactic nuclei (hereafter \\ppagn) have been found since the 1980s (e.g. Heckman et al. 1981, 1984; Keel 1985; Whittle 1985a,b,c), but were mainly interpreted as an indicator of outflows from active nuclei or rotating narrow line regions (e.g. Greene \\& Ho 2005). Since the suggestion that double-peaked profiles could be produced by dual AGN (Zhou et al. 2004; Wang et al. 2009), much attention has been recently paid to search for individual objects and study statistical properties of \\ppagn\\ samples in light of several fundamental issues related with evolution of galaxies and AGN (for a summary of previous works on this subject see Table 1). There are 8 individual objects with double-peaked narrow emission-line profiles in total listed in Table 1, and three samples that made systematic searches for \\ppagn\\ from SDSS DR7 (Wang et al. 2009; Liu et al. 2010a; Smith et al. 2010; hereafter W09, L10 and S10, respectively). W09, L10 and S10 select their samples by employing different selection criteria, including restricting equivalent width (EW), redshift, signal-to-noise ratio (S/N), narrow emission line (NEL) flux ratio etc. W09 and L10 select 190 \\ppagn\\ in total from SDSS galaxy sample using Kewley et al. (2001) criterion in the BPT diagram (Baldwin et al. 1981). This paper applies new selection scheme and criteria to hunt double-peaked narrow emission-line galaxies (\\ppgalaxies) from $\\sim 920,000$ galaxies.  Double-peaked narrow line profiles can be caused by different mechanisms discussed subsequently. First, bi-polar outflows driven by starbursts or AGN radiation pressure have been suggested for many years, especially in several individual objects, such as NGC 1068 (Crenshaw \\& Kraemer 2000; Das et al. 2006, 2007), NGC 4151 (Hunchings et al. 1998; Crenshaw et al. 2000; Das et al. 2005), Mrk 3 (Ruiz et al. 2001; Crenshaw et al. 2010a), Mrk 78 (Fisher et al. 2011), and Mrk 573 (Schlesinger et al. 2009; Fisher et al. 2010). In particular, outflows can produce double-peaked (e.g. Mrk 78 see Fisher et al. 2011) or multiple-peaked emission-line profiles (e.g. NGC 1068, see also Crenshaw et al. 2010b) depending on the projection of the outflows in the line of sight. However, W09 found an anti-correlation between the flux ($F_{\\rm r}/F_{\\rm b}$) and shift ($\\Delta \\lambda_{\\rm r}/\\Delta \\lambda_{\\rm b}$) ratios of the double peaks, where $F$ and $\\Delta \\lambda$ are the fluxes and shifts of red and blue components, which has been confirmed by Liu et al. (2010a, see their Figure 6{\\it e}), Smith et al. (2010, see their Figure 5) and Fu et al. (2011, see their Figure 5). This new statistical relation strongly suggests co-rotating dual AGN model. Unless the bi-polar outflows with very asymmetric mass rates are under momentum conservation, the anti-correlation would not favor the presence of bi-polar outflows (W09)\\footnote{Here we assume that the outflows are optically thin. The radiation from them is proportional to the mass of the emitting gas expected.}. Second, rotating disk-like narrow line regions (NLRs) with circular orbits are {\\it not} able to produce the anti-correlation because a large fraction of \\ppagn\\ have $F_{\\rm r}/F_{\\rm b}\\neq 1$, unless the rotating NLRs have complicated elliptical orbits (or inhomogeneous mass distribution over the rotating disk) and orientations to observers. Third, the origin of double-peaked profiles could be interpreted by dual AGN which will be in agreement with the anti-correlation (W09; L10; Smith et al. 2010; Rosario et al. 2011) as identified by some works (e.g. Fu, et al. 2011; Shen, et al. 2011) listed in Table 1. In order to identify the nature of these $pp$-AGN, near infrared (NIR) images and optical long slit spectroscopy have shown that about 50\\% type I $pp$-AGN (Fu et al. 2011), 10-20\\% type II $pp$-AGN (Liu et al. 2010b; Shen et al. 2011; Fu et al. 2011, 2012), or more than $50\\%$ \\ppagn\\ (Comerford et al. 2011) and 50\\% double-peaked quasars (Rosario et al. 2011) have two cores. The nature of the remaining objects is open (a single core: bi-polar outflows or rotating NLR; or spatially unresolved dual-cores).  Double-peaked narrow emission-line galaxies play an important role in studies of the dynamics of AGN NLRs and merging galaxies. The formation of NLRs is insufficiently understood even though they are spatially resolved in some individual objects (Bennert et al. 2002; Schmitt et al. 2003; Bennert et al. 2006a,b). Jet-induced outflows are responsible for NLR formation in some AGN (Bicknell et al. 1998) in light of individual objects, such as Mrk 3 (Capetti et al. 1999), and Mrk 78 (Whittle \\& Wilson 2004). More clear evidence for the jet-clouds interaction in NGC 4051 has been revealed by {\\it Chandra} (Wang et al. 2011). Cold clouds are formed by the cocoons produced through the Kelvin-Helmholtz instability of the relativistic jet propagating in the interstellar medium (ISM) (Steffen et al. 1997). NLRs driven by the intermediate scale of outflows or jets can cause the appearance of double-peaked profiles (Rosario et al. 2008, 2010). This is supported by the fact that there is a correlation between the jet power and width of NELs (Wilson \\& Willis 1980) and powerful, compact linear radio sources have anomalously broad \\oiii\\ lines usually (Whittle 1992). It remains to be systematically researched on NLR dynamics driven by jet-induced outflows in a large and homogeneous sample. Furthermore, feedback to starburst through AGN outflows is presumed to play a key role in regulating bulge growth, but it still needs to be evidenced by observations of a large \\ppgalaxies\\ sample driven by bi-polar outflows. On the other hand, Greene \\& Ho (2005) conducted a systematic investigation of NLR dynamics from a large SDSS sample. They draw a conclusion that the NLR is mainly governed by the potential of bulges though there are some evidence for rotation-supported NLR in light of Integral Field Unit (IFU) rotation curves (Vega Beltr\\'an et al. 2001; Dumas et al. 2007; Stoklasov\\'a et al. 2009). Clearly, those \\ppagn with double-peaked NELs produced by a rotation-supported NLR would be used to justify if the rotation-supported NLR follows the dynamics of bulges (Dumas et al. 2007; Ho 2009) or the rotation of partial disks of galaxies, even provide an opportunity to study the relative orientation of AGN and host galaxies (Shen et al. 2010; Lagos et al. 2011).  Merging galaxies enter a special phase in which double cores exist after their galactic disks merge within the framework of the theory of hierarchical galaxy formation and evolution (e.g. Longair 2008; Yu et al. 2011). As illustrated by SDSS images (e.g. Figure 2 in Darg et al. 2010), there is a sequence of galaxy mergers from tidal interaction between two galaxies with a separation of tens of kpc to an ``approaching post-merger\" stage containing dual cores with a $\\sim$kpc separation. If hierarchical growth of galaxies works, are there AGN-starburst cores, or galaxies with dual starbursts? How often? Do they appear as \\ppgalaxies? If so, these \\ppgalaxies\\ could open a new clue to understanding the issues related to triggering starbursts and black hole activity, and help to understand the details of mergers. Second, dual-cored galaxies are important for the subject of black hole binaries searched through shifts between broad lines and narrow lines (Tsalmantza et al. 2011, Eracleous et al. 2011) since dual-cored galaxies are progenitors of binary black holes. With the dual-cored galaxies, we will have an entire scenario of merging process. These motivate us to systematically search for \\ppgalaxies\\ to look into statistical properties.  In this paper, we make an attempt to establish the largest sample of \\ppgalaxies\\ through a systematical search from SDSS DR7 spectra. As a by-product, we also build a sample of asymmetric or top-flat narrow line (asym-NEL) galaxies, some of which could be candidates of \\ppgalaxies. Section 2 describes the followed procedures, selection criteria and the final samples of this work. We systematically inspect images of the \\ppgalaxies\\ and asym-NEL galaxies in Section 3 and find 54 of them with dual cores. In Section 4, we present the basic properties of the sample. We draw conclusions in Section 5. We use a standard cosmology with $H_0=71~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_M=0.27$ and $\\Omega_{\\Lambda}=0.73$ in this paper. ",
        "conclusions": "We have developed a pipeline to automatically select candidates of double-peaked galaxies of spectra. Combining with visual inspection of spectra, we set up the largest sample of double-peaked galaxies to date, which is composed of double-peaked narrow emission-line galaxies (3,030). We also build a sample of galaxies (12,582) with asymmetric and top-flat profiles from the SDSS DR7 galaxy sample. The \\ppgalaxies\\ and the asym-NEL galaxies show very similar statistical properties, but velocity separations are significantly different. We find that double-peaked components are divided into 5 kinds, namely, type I \\ppagn, 2-type II \\ppagn, (AGN+SF) \\ppgalaxies, 2-SF and unknown-type objects depending on the Balmer line width and the Baldwin-Phillips-Terlevich diagram. The \\ppgalaxies\\ sample covers: 81 type I \\ppagn, 837 2-type II \\ppagn, 400 type II+SF \\ppgalaxies, 708 2-SF and 1,004 unknown-type objects. We visually inspect the SDSS images of the sample and find 54 galaxies with dual-cores, of which 14 2-type II \\ppagn, 8 2-SF, 12 type II+SF \\ppgalaxies\\, and 20 unknown-type objects.  We would draw the following conclusions in light of the present sample: \\begin{itemize} \\item The fractions of \\ppgalaxies\\ and asym-NEL galaxies to emission-line galaxies are about $1.0\\%$ and $3.6\\%$, respectively, and \\ppagn\\ (including type I, 2-type II, type II +SF \\ppgalaxies) and asym-NEL AGN are about $0.4\\%$ and $1.5\\%$, respectively.  \\item From the two samples, we find 54 dual cored galaxies with projected separation smaller than $3^{\\pp}$ and 255 with that larger than $3^{\\pp}$. Observations with higher spatial resolution or in X-ray or radio bands are needed to identify dual-cores from images of rest objects of the present samples, in particular, the sub-samples of type I and AGN+SF \\ppgalaxies. \\end{itemize}  Future observations with for example the {\\it LAMOST/Guoshoujing} (a 4-m telescope with 4,000 fibers) through images and spectra of the two components are needed to understand the nature of the $pp$-galaxies. The large sample can be used to study the dynamical processes of merging galaxies, the triggering mechanisms of AGN activity, outflows and rotation-supported disk of narrow line regions."
    },
    "1208/1208.2957_arXiv.txt": {
        "abstract": "Several properties of the Solar System, including the wide radial spacing and orbital eccentricities of giant planets, can be explained if the early Solar System evolved through a dynamical instability followed by migration of planets in the planetesimal disk. Here we report the results of a statistical study, in which we performed nearly $10^4$ numerical simulations of planetary instability starting from hundreds of different initial conditions. We found that the dynamical evolution is typically too violent, if Jupiter and Saturn start in the 3:2 resonance, leading to ejection of at least one ice giant from the Solar System. Planet ejection can be avoided if the mass of the transplanetary disk of planetesimals was large ($M_{\\rm disk}\\gtrsim50$ M$_{\\rm Earth}$), but we found that a massive disk would lead to excessive dynamical damping (e.g., final $e_{\\rm 55} \\lesssim0.01$ compared to present $e_{\\rm 55}=0.044$, where $e_{55}$ is the amplitude of the fifth eccentric mode in the Jupiter's orbit), and to smooth migration that violates constraints from the survival of the terrestrial planets. Better results were obtained when the Solar System was assumed to have five giant planets initially and one ice giant, with the mass comparable to that of Uranus and Neptune, was ejected into interstellar space by Jupiter. The best results were obtained when the ejected planet was placed into the external 3:2 or 4:3 resonance with Saturn and $M_{\\rm disk} \\simeq 20$ M$_{\\rm Earth}$. The range of possible outcomes is rather broad in this case, indicating that the present Solar System is neither a typical nor expected result for a given initial state, and occurs, in best cases, with only a $\\simeq$5\\% probability (as defined by the success criteria described in the main text). The case with six giant planets shows interesting dynamics but does offer significant advantages relative to the five planet case. ",
        "introduction": "As giant planets radially migrate in the protoplanetary nebula they should commonly be drawn into compact systems, in which the pairs of neighbor planets are locked in the orbital resonances (Kley 2000, Masset \\& Snellgrove 2001). The resonant planetary systems emerging from the protoplanetary disks can become dynamically unstable after the gas disappears, leading to a phase when planets scatter each other. This model can help to explain the observed resonant exoplanets (e.g., Gliese 876; Marcy et al. 2001), commonly large exoplanet eccentricities (Weidenschilling \\& Marzari 1996, Rasio \\& Ford 1996), and microlensing data that show evidence for a large number of planets that are free-floating in interstellar space (Sumi et al. 2011; but see Veras \\& Raymond 2012).  The Solar System, with the widely spaced and nearly circular orbits of the giant planets, bears little resemblance to the bulk of known exoplanets. Yet, if our understanding of physics of planet--gas-disk interaction is correct, it is likely that the young Solar System followed the evolutionary path outlined above. Jupiter and Saturn, for example, were most likely trapped in the 3:2 resonance (Masset \\& Snellgrove 2001, Morbidelli \\& Crida 2007, Pierens \\& Nelson 2008, Pierens \\& Raymond 2011, Walsh et al. 2011), defined as $P_{\\rm Sat}/P_{\\rm Jup}=1.5$, where $P_{\\rm Jup}$ and $P_{\\rm Sat}$ are the orbital periods of Jupiter and Saturn (this ratio is 2.49 today).  To stretch to its current configuration, the outer Solar System most likely underwent a dynamical instability during which Uranus and Neptune were scattered off of Jupiter and Saturn and acquired eccentric orbits (Thommes et al. 1999, Tsiganis et al. 2005, Morbidelli et al. 2007, Levison et al. 2011). The orbits were subsequently stabilized (and circularized) by damping the excess orbital energy into a massive disk of planetesimals located beyond the orbit of the outermost planet (hereafter transplanetary disk), whose remains survived to this time in the Kuiper belt. Finally, as evidenced by dynamical structures observed in the Kuiper belt, planets radially migrated to their current orbits by scattering planetesimals (Malhotra et al. 1995, Gomes et al. 2004, Levison et al. 2008).  Several versions of the Solar System instability were proposed. Thommes et al. (1999), motivated by the apparent inability of the existing formation models to accrete Uranus and Neptune at their present locations, proposed that Uranus and Neptune formed in the Jupiter-Saturn zone, and were scattered outwards when Jupiter, and perhaps Saturn, accreted nebular gas. This work represented an important paradigm shift in studies of the Solar System formation. Inspired by this work, Levison et al. (2001) suggested that the instability and subsequent dispersal of the planetesimal disk, if appropriately delayed (or due to the late formation of Uranus and Neptune; Wetherill 1975), could explain the Late Heavy Bombardment (LHB) of the Moon. The LHB was a spike in the bombardment rate some 4 Gyr ago suggested by the clustering of ages of several lunar basins. The nature of the LHB, however, is still being debated (see Hartmann et al. 2000 and Chapman et al. 2007 for reviews).  In the {\\it Nice} version of the instability (Tsiganis et al. 2005), Uranus and Neptune were assumed to have formed just outside the orbit of Saturn, while Jupiter and Saturn were assumed to have initial orbits (where ``initial'' means at the time when the protoplanetary nebula dispersed), such that $P_{\\rm Sat}/P_{\\rm Jup}<2$. The instability was triggered when Jupiter and Saturn, radially drifting by scattering planetesimals, approached and crossed the 2:1 resonance ($P_{\\rm Sat}/P_{\\rm Jup}=2$). The subsequent evolution was similar to that found in the Thommes et al. model, but led to a better final orbital configuration of the planets (assuming that the disk contained $\\sim20$-50 M$_{\\rm Earth}$ and was truncated at 30-35 AU).  The {\\it Nice} model is compelling because it explains the separations, eccentricities, and inclinations of the outer planets (Tsiganis et al. 2005), and many properties of the populations of small bodies (Morbidelli et al. 2005, Nesvorn\\'y et al. 2007, Levison et al. 2008, 2009, Nesvorn\\'y \\& Vokrouhlick\\'y 2009). It is currently the only migration model that is consistent with the current dynamical structure of the terrestrial planets (Brasser et al. 2009) and the main asteroid belt (Morbidelli et al. 2010). Moreover, the {\\it Nice} model can also reproduce the magnitude and duration of the LHB (Gomes et al. 2005, Bottke et al. 2012).  Two potentially problematic issues of the original {\\it Nice} model (hereafter ONM) were addressed by Morbidelli et al. (2007) and Levison et al. (2011). Morbidelli et al. pointed out that the initial planetary orbits in the ONM were chosen without the proper regard to the previous stage of evolution during which the giant planets interacted with the protoplanetary nebula. As shown by hydrodynamical studies (Masset \\& Snellgrove 2001, Morbidelli \\& Crida 2007, Pierens \\& Nelson 2008, Pierens \\& Raymond 2011), this interaction most likely led to the convergent migration of planets, and their trapping in orbital resonances. Morbidelli et al. (2007) studied the dynamical instability starting from the resonant planetary systems and showed that the orbit evolution of planets was similar to that of the ONM.  To produce the LHB, the instability needs to occur late relative to the dispersal of the protoplanetary nebula. A protoplanetary nebula typically disperses in $\\sim$3-10 Myr after the birth of the star (Haisch et al. 2001), which for the Sun was probably contemporary to the formation of the first Solar System solids, 4.568 Gyr ago (Bouvier et al. 2007, Burkhardt et al. 2008). The onset of the LHB, traditionally considered to be 3.9-4.0 Gyr ago (Ryder 2002), has been recently revised to 4.1-4.2 Gyr ago (Bottke et al. 2012). This means that the instability had to occur with a delay of 350-650 Myr, with the exact value depending on the actual time of the LHB event.\\footnote{Here we opted for including the full range of previous and new estimates of the delay, because it is not 100\\% guaranteed that the new estimate is correct. Imposing the 350-450 My delay, instead of 350-650 My, would not make much of a difference, because most planetary systems that are stable over 450 My are also stable over 650 My.} Such a late instability can be triggered in the ONM if the inner edge of the planetesimal disk was close, but not too close, to the outer ice giant. If the edge had been too close, the instability would have happened too early to be related to the LHB. If the edge had been too far, the instability would have happened too late or would not have happened at all, because the planetesimals would had stayed radially confined and not evolved onto planet-crossing orbits (where they could influence planets by short-range interactions).  To avoid the need for fine tuning of the inner disk's edge in the ONM, Levison et al. (2011) proposed that the late instability was caused by the {\\it long-range} interactions between planets and the radially-confined distant planetesimal disk. The long-term interactions arise when gravitational scattering between disk's planetesimals is taken into account. As a result of these interactions, the planets and planetesimal disk slowly exchanged the energy and angular momentum until, after hundreds of Myr of small changes, the resonant locks between planets were broken and the instability reigned supreme. Note that the instability trigger in Levison et al. (broken resonant locks) is fundamentally different from that of the ONM (major resonance crossing during migration).  Here we report the results of a new statistical study of the Solar System instability. Sections 2-4 explain our method and constraints. The results are presented in Sect. 5. We discuss the plausible initial configurations of planets after the gas nebula dispersal, mass of the planetesimal disk, and effect of different trigger mechanisms on the results. The first steps in this direction were taken by Batygin \\& Brown (2010), who found that some initial resonant configurations may work while others probably do not.  Following Nesvorn\\'y (2011) and Batygin et al. (2012) we also consider cases in which the young Solar System had extra ice giants initially and lost them during the scattering phase (these works and their relation to the present paper will be discussed in more detail below; see, e.g., Sect. 6). We extend the previous studies of the Solar System instability by: (1) exploring a wide range of initial parameters (new resonant chains, six planets, etc.), (2) improving the statistics with up to 100 simulations per case, (3) considering different trigger mechanisms (including the one recently proposed by Levison et al. 2011), and (4) applying strict criteria of success/failure to all studied cases (see Sect. 3 and 4). Our conclusions are given in Sect. 7. ",
        "conclusions": "Recent studies suggest that Jupiter and Saturn formed and migrated in the protoplanetary gas disk to reach a mutual resonance, most likely the 3:2 resonance, where the Saturn's orbital period was 3/2 longer than that of Jupiter. After the gas disk's dispersal, the orbits of Jupiter and Saturn must have evolved in some way to eventually arrive to the current orbits with the orbital period ratio of 2.49. This can most easily be achieved, considering constraints from the terrestrial planets and $e_{55}$, if Jupiter and Saturn scattered off of Uranus or Neptune, or a planet with mass similar to that of Uranus or Neptune.  We performed $N$-body integrations of the scattering phase between the Solar System's giant planets, including cases where one or two extra ice giants were assumed to have formed in the outer Solar System and ejected into interstellar space during instability. We found that the initially compact resonant configurations and low masses of the planetesimal disk ($M_{\\rm disk}<50$~M$_{\\rm Earth}$) typically lead to violent instabilities and planet ejection. On the other hand, the initial states with orbits that are more radially spread (e.g., Jupiter and Saturn in the 2:1 resonance) and larger $M_{\\rm disk}$ result in smooth migration of the planetary orbits that leads to incorrectly low $e_{55}$ and excitation of the terrestrial planet orbits. Finding the sweet spot between these two extremes is difficult.  Some of the statistically best results were obtained when assuming that the Solar System initially had five giant planets and one ice giant, with the mass comparable to that of Uranus and Neptune, was ejected to interstellar space by Jupiter. The best results were obtained when the fifth planet was assumed to have the mass similar to Uranus/Neptune, was placed on an orbit just exterior to Saturn's (3:2 and 4:3 resonances work best), and the orbits of Uranus and Neptune migrated into the planetesimal disk before the onset of planetary scattering. This mode of instability is favored for several reasons, as described below.  As planetesimals are scattered by Uranus and Neptune and evolve into the Jupiter/Saturn region, Jupiter, Sa\\-turn and the fifth planet undergo divergent migration. This triggers an instability during which the fifth planet suffers close encounters with all planets and is eventually ejected from the Solar System by Jupiter. Uranus and Neptune generally survive the scattering phase, because their orbits migrated outward during the previous stage and opened a protective gap between them and the gas giants. This mode of instability produces just the right kind of Jupiter's semimajor axis evolution -known as jumping Jupiter- that is required from the terrestrial planet constraint.  Moreover, $e_{55}$, excited by the fifth planet ejection, is not damped to incorrectly low values by secular friction from the planetesimal disk, because the planetesimal disk had been disrupted by Uranus and Neptune before the excitation event. The low mass of the planetesimal disk at the time of planet scattering also leads to only a brief migration phase of Jupiter and Saturn after the scattering phase, and prevents $P_{\\rm Sat}/P_{\\rm Jup}$ from evolving beyond its current value. The excessive residual migration of Jupiter and Saturn was a problem in most other cases investigated here.  The mode of instability with early migration of Uranus and Neptune is problematic, however, because the inner edge of the planetesimal disk may need to be fine-tuned to generate the delay between the formation of the Solar System and the LHB. In addition, the range of possible outcomes is rather broad, indicating that the present Solar System is neither typical nor an expected result, and occurs, in best cases, with only a $\\simeq$5\\% probability (as defined by simultaneous matching of all four criteria defined in Sect. 3). In $\\simeq$95\\% of cases, the simulations ended up failing at least one of our constraints. This may seem unsatisfactory, but given the issues discussed in Sect. 4, the fact that our criteria are relatively strict, and because the instability-free model does not work at all,\\footnote{The instability-free model with four planets can be easily tuned to satisfy B, but not C and D, and constraints from the small body populations (e.g., Walsh \\& Morbidelli 2011, Dawson \\& Murray-Clay 2012).} these findings should be seen in a positive light. An important follow-up of this work will be to consider additional constraints from the small body populations (e.g., the dynamical structure of the Kuiper belt), and see whether the successful simulations identified here will also match those constraints.  The case with six giant planets is also interesting in that the instability occurs in two steps, corresponding to the ejection of the two planets. The best results were obtained in this case when the two ejected planets were given similar masses (about half the mass of Neptune) and were placed between the orbits of Saturn and the inner surviving ice giant. As expected, the six-planet case leads to a larger variety of results than the five-planet case. The probability of ending the six-planet simulation with the present properties of the solar system is therefore lower than in the five-planet case. Still, our six-planet results are fundamentally better than those obtained in the four-planet case, where the differences were systematic (e.g., $e_{\\rm 55}$ never large enough), and the success rate was below the resolution limit of our study."
    },
    "1208/1208.0449_arXiv.txt": {
        "abstract": "{A model of unification of dark matter and dark energy based on the modeling of the speed of sound as a function of the parameter of the equation of state is introduced. It is found that the model in which the speed of sound depends on the power of the parameter of the equation of state, $c_s^2=\\alpha (-w)^{\\gamma}$, contains the generalized Chaplygin gas models as its subclass. An effective scalar field description of the model is obtained in a parametric form which in some cases can be translated into a closed form solution for the scalar field potential. A constraint on model parameters is obtained using the observational data on the Hubble parameter at different redshifts.}  \\vspace{2cm} ",
        "introduction": "The expansion of the universe is one of the most fascinating phenomena that science has encountered so far. It has served as a rich source of information on the nature and the composition of the universe. One of recently established astonishing features of cosmic expansion is that it is currently undergoing a phase of acceleration \\cite{R98,P99,S03,T04}. The source of this acceleration has not yet been unambiguously identified, although many proposals for its nature have been put forward (see \\cite{Copeland,Saridakis} and references therein). The existence of dark energy, a mysterious component with negative pressure is still the most serious candidate. Another dark component of the universe, dark matter, seems to leave its imprint on astrophysical and cosmological scales, ranging from galaxies to galactic clusters and large scale structure in the universe.  The idea that both dark matter and dark energy are actually the manifestations of a single dark component is both natural and appealing. It appeared early in the literature and its the most acclaimed representative is probably the Chaplygin gas \\cite{K01} as a model of unification of dark matter and dark energy \\cite{Bi02}. The class of unifying DM-DE models is often referred to as quartessence \\cite{Makler,Reis}. This phenomenologically introduced model can be motivated from string theory \\cite{Bi02,O00,R00}. Its usually studied extension, the generalized Chaplygin gas model, was  first introduced in \\cite{B02}. The agreement of the Chaplygin gas and its extensions with observations has been extensively tested, including the analyses with the supernova Ia data \\cite{FGS02,C08}, CMB \\cite{AFBC03}, observable Hubble parameter data \\cite{Wu,F11} and large scale structure observations \\cite{ST04,B03,F08,GK08,F10}, including the nonlinear evolution in structure formation \\cite{Bilic2004} and gravastar formation \\cite{Bilic2006Grav}. Different data can be combined to produce tighter parameter constrains such as in \\cite{P10,BD03,CV09,LWY09}. Strong constraints on the generalized Chaplygin gas have been obtained that question its viability as a cosmological model distinguishable from the $\\Lambda$CDM model. In order to better accommodate observational constrains, various unified models based on Chaplygin gas have been proposed, such as the modified Chaplygin gas model \\cite{barrow,benaoum}, recently reviewed  and constrained in \\cite{L08,D11,LXWL11} or hybrid Chaplygin gas leading to transient acceleration \\cite{bilic2005trans}. The fact that perfect fluid model can be fully described by defining speed of sound equation has been used in \\cite{X11}. The idea of DM-DE unification with non-canonical scalar fields has been recently studied in \\cite{BB07,DFP12}. An interesting model called {\\em Dusty Dark Energy}, recently introduced in \\cite{Vikman}, achieves the DM-DE unification in the formalism of the $\\lambda \\varphi${\\em -fluid}, resulting in the zero speed of sound and one scalar degree of freedom. Other approaches to models of DM-DE unification that avoid the speed of sound problem are purely kinetic k-essence models \\cite{S04,B07,Chimento} and tachyon models \\cite{Padmanabhan,bilic2009tach}.  The structure of the paper is the following. After the introduction presented in this section, in the second section a general class of barotropic fluid models defined by the function $c_s^2$ is discussed. In the third section the model with the constant speed of sound is studied and in the fourth section the principal model of the paper, defined by $c_s^2(w)=\\alpha (-w)^{\\gamma}$ is introduced. The fifth section is focused on an effective representation of the model in terms of a minimally coupled scalar field. In the sixth section the comparison of the model prediction against the observational data on the Hubble parameter at different redshifts is made and the constraints on the model parameters are presented. The seventh section closes the paper with the discussion and conclusions. The Appendix outlines an approach in which the solution with the piecewise constant speed of sound is used as an approximation of the dynamics of a fluid with a general dependence of $c_s^2$ on $w$. ",
        "conclusions": "The analysis presented in the preceding section serves primarily as an orientation regarding the allowed part of the parametric $(\\alpha,\\gamma)$ space. A more complete analysis should consider the growth of inhomogeneities and make the comparison of the model predictions with other observational data, such as the data on the matter power spectrum, supernovae of the type Ia and cosmic microwave background. Although preliminary analyses in this direction have already been made \\cite{CaplarDipl}, a detailed analysis of observational data is left for future work \\cite{Prep}.  From Figs \\ref{fig:chi2} and \\ref{fig:chi2H0} it is evident that for larger values of the exponent $\\gamma$, a wider interval of the coefficient $\\alpha$ is allowed. A similar conclusion follows from the preliminary analysis of the matter power spectrum \\cite{CaplarDipl}. This feature can be understood on qualitative grounds using the following argumentation. For large values of $\\gamma$, as long as the EOS parameter is not far from $w=0$, the speed of sound of the dark component remains small and suppressed by the large exponent $\\gamma$. This fact prevents the formation of the sound horizon and its adverse effects on structure formation.  The generalized model of unification of dark matter and dark energy introduced in this paper opens a novel perspective on Chaplygin gas and its modifications and generalizations. The model defined by the relation $c_s^2=\\alpha (-w)^{\\gamma}$ encompasses both the Chaplygin gas (for $\\gamma=1$, $\\alpha=1$) and the generalized Chaplygin gas (for $\\gamma=1$) as a specific subclass. It could therefore serve as a wider framework for the analysis how much these models need to be extended to satisfy the constraint from the observational data. From the modeling perspective, the crucial element of our model is the relation between the speed of sound, $c_S^2$, and the parameter of the EOS, $w$. This relation connects the quantity governing the growth of inhomogeneities with the quantity determining the global evolution of the universe. This feature might allow easier and more direct transformation of the phenomenological knowledge acquired from the data into workable models of the dark sector. In particular, an adaptive model assuming piecewise constant values of the speed of sound in consecutive redshift intervals is presented in the Appendix.  A particular challenge for the future research is finding a microscopic explanation of the dependence $c_S^2(w)$. Here the corresponding microscopic models for the Chaplygin gas might serve as a good starting point. In particular, the method of representing the evolution in terms of minimally coupled scalar fields, presented in section \\ref{sec:scalar} could provide useful information on such microscopic models.   \\vspace{0.5cm}   {\\bf Acknowledgements.} H. \\v{S}. acknowledges useful discussions with D. Huterer in the early phase of this work. The authors would like to thank N. Bili\\'{c} for useful comments on the manuscript. This work was supported by the Ministry of Education, Science and Sports of the Republic of Croatia under the contract No. 098-0982930-2864."
    },
    "1208/1208.0480_arXiv.txt": {
        "abstract": "{ Rapidly oscillating Ap (roAp) stars have rarely been found in binary or higher order multiple systems. This might have implications for their origin. } { We intend to study the multiplicity of this type of chemically peculiar stars, looking for visual companions in the range of angular separation between 0\\farcs05 and 8\\arcsec{}. } { We carried out a survey of 28 roAp stars using diffraction-limited near-infrared imaging with NAOS-CONICA at the VLT. Additionally, we observed three non-oscillating magnetic Ap stars. } { We detected a total of six companion candidates with low chance projection probabilities. Four of these are new detections, the other two are confirmations. An additional 39 companion candidates are very likely chance projections. We also found one binary system among the non-oscillating magnetic Ap stars. The detected companion candidates have apparent K magnitudes between 6\\fm8 and 19\\fm5 and angular separations ranging from 0\\farcs23 to 8\\farcs9, corresponding to linear projected separations of 30--2400\\,AU. } { While our study confirms that roAp stars are indeed not very often members of binary or multiple systems, we have found four new companion candidates that are likely physical companions. A confirmation of their status will help understanding the origin of the roAp stars. } ",
        "introduction": "\\label{sect:intro}    Rapidly oscillating Ap (roAp) stars are ideal targets for asteroseismology. By comparing the observed frequency spectrum with the asymptotic pulsation theory, it is possible to specify their rotation period, their temperature, luminosity, radius, mass, their atmospheric structure, their evolutionary status, and the geometry of their magnetic field. More than forty roAp stars are known at present, with effective temperatures between 6400\\,K and 8400\\,K. Kurtz et al.\\ (\\cite{Kurtz2006}) list 35 roAp stars; more can be found in individual papers since then. They pulsate in high-overtone, low-degree, nonradial $p$-modes, with periods in the range from 5.6 to 21\\,min and typical amplitudes of a few millimagnitudes (e.g.\\ Kurtz \\cite{Kurtz1982}).  The roAp phenomenon is confined to a well-defined region of the Str\\\"omgren photometry parameter space (Martinez \\cite{Martinez1993}). However, this region also contains other Ap stars, in which no pulsation could be detected, despite thorough searches. These apparently constant Ap stars (non-oscillating Ap stars, or noAp stars) appear remarkably similar to the roAp stars in many respects (color indices, abundances, magnetic fields).  A decade ago, after the completion of a study of the kinematical properties of rapidly oscillating Ap and noAp stars, Hubrig et al.\\ (\\cite{Hubrig2000}) realized that none of the roAp stars is known to be a spectroscopic binary (SB). They obtained several radial velocity measurements for a majority of the roAp stars, but found no evidence for variations in any of these stars. The situation is quite different for other Ap stars co-existing in the same region of the H-R diagram. In contrast with roAp stars, many noAp stars for which radial velocity data exist are known as spectroscopic binaries or show radial velocity variations. The interpretation of this difference and its significance for the understanding of the origin of the pulsations in roAp stars is not clear so far. That until now no roAp star is known to be a spectroscopic binary is also in direct contrast to the situation for other types of pulsating variables (e.g., $\\beta$\\,Cep stars, $\\delta$ Sct stars, or classical Cepheids), which are frequently found in SB systems.  Neither theoretically nor observationally is our present knowledge sufficient to decide confidently whether tidal interaction in binaries may reduce the amplitude of or inhibit the pulsations in cool Ap stars. To establish this, a necessary condition would be to show that essentially all noAp stars are close binaries, or, alternatively, to investigate whether all roAp stars are single stars or wide visual binaries.  Our search in the literature and catalogs for known double or multiple systems among roAp stars revealed that the three roAp stars $\\gamma$\\,Equ (HD\\,201601), $\\beta$\\,CrB, and $\\alpha$\\,Cir, which are the brightest and best studied stars in the group of roAp stars, do have optical or physical companions. Stelzer et al.\\ (\\cite{Stelzer2011}) combined data from the Washington Double Star Catalog (Mason et al.\\ \\cite{Mason2001}) and from {\\em Hipparcos} with the measurement reported in this paper and derived a preliminary orbit for $\\gamma$\\,Equ. They estimated the mass of the companion to be 0.6$\\pm$0.4\\,$M_{\\sun}$. The companion in the $\\beta$\\,CrB system at a separation of 0\\farcs3 was frequently observed by speckle interferometry in the eighties. The companion in the system $\\alpha$\\,Cir is of spectral type K5V (Eggleton \\& Tokovinin \\cite{EggletonTokovinin2008}). All other roAp stars are considerably fainter than $\\gamma$\\,Equ, $\\beta$\\,CrB, and $\\alpha$\\,Cir, and therefore have not been intensely studied with high resolution imaging instruments. Still, companions are reported for HD\\,99563 and HR\\,3831 (e.g.\\ Dommanget \\& Nys \\cite{DommangetNys2002}).  The interpretation of a difference in duplicity between roAp and noAp stars and its meaning for the understanding of the origin of the pulsations in roAp stars is far from obvious. Even if the different internal structure of roAp stars is the reason for their pulsations, it is difficult to understand why no roAp star was found to be in a close binary.   In the following we report the results of our multiplicity study of this class of objects using NACO K-band imaging. ",
        "conclusions": "\\label{sect:discussion}  \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth, angle=0]{sep_histo.eps} \\caption{ Distribution of the projected separations of the studied systems with roAp primaries. For this figure, we have used all companion candidates where a parallax was available for the roAp star. The shaded region shows the three objects where the chance projection probability is smaller than 1\\%. } \\label{fig:histo} \\end{figure}  \\begin{table} \\centering \\caption{ Overview of the known multiplicity of the objects studied in this article. } \\label{tab:multi} \\begin{tabular}{rccc} \\hline \\hline \\multicolumn{1}{c}{HD} & \\multicolumn{1}{c}{SB1} & \\multicolumn{1}{c}{Astrometric} & \\multicolumn{1}{c}{Comment} \\\\ \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{or Visual} & \\multicolumn{1}{c}{} \\\\ \\hline 6532   &    &   & \\\\ 9289   &    & 1 & ND \\\\ 12932  &    & 1 & ND \\\\ 19918  &    &   & \\\\ 24712  &    &   & \\\\ 42659  &    &   & \\\\ 86181  &    &   & \\\\ 99563  &    & 1 & CD \\\\ 101065 &    &   & \\\\ 116114 & X? & 1 & NR \\\\ 119027 &    &   & \\\\ 122970 &    &   & \\\\ 134214 &    &   & \\\\ 137949 &    &   & \\\\ 150562 &    &   & \\\\ 154708 &    & 1 & ND \\\\ 161459 &    &   & \\\\ 166473 &    &   & \\\\ 176232 &    &   & \\\\ 185256 &    &   & \\\\ 190290 &    &   & \\\\ 193756 &    &   & \\\\ 196470 &    &   & \\\\ 201601 &    & 3 & CD+NR+NR\\\\ 203932 &    & 1 & ND \\\\ 213637 &    &   & \\\\ 217522 &    &   & \\\\ 218495 &    &   & \\\\ \\hline \\multicolumn{3}{c}{Other A stars} \\\\ \\hline 40711  & X? &   & \\\\ 55719  & X  & 1 & CD \\\\ 59435  & X? &   & \\\\ \\hline \\end{tabular} \\tablefoot{ An X in the second column denotes an SB1 system; the ? indicates objects where there are only hints for an SB1 system. In the last column we point out the new detections (ND), the confirmed detections (CD), and the not reachable companions (NR). The not reachable companions are either potentially very narrow (HD\\,116114) or outside of our field-of-view (HD\\,201601). } \\end{table}  Here, we announce the detection of 45 companion candidates in 13 of the observed 28 roAp stars, with 39 of these very likely being chance projections. They have K magnitudes between 6\\fm8 and 19\\fm5 and angular separations ranging from 0\\farcs23 to 8\\farcs9, corresponding to linear projected separations of 30--2400\\,AU. For the eight companion candidates in  potential binary systems, three have a high probability to be chance projections. In the systems with more than one companion candidate, all companion candidates but HD\\,154708AC are very likely chance projections. All companion candidates, except for the companion candidates of HD\\,99563 and HD\\,201601, were detected by us for the first time. The companion candidate for the noAp star HD\\,55719 was also known before. In Fig.~\\ref{fig:histo} we show the distribution of the projected linear separations for the studied multiple systems with roAp primaries. The distribution is missing the 26 companion candidates where no parallax information is available for the roAp star. For only three companion candidates with low probability to be a chance projection exist parallax data. These are the three objects with a projected linear separation below 500\\,AU.  In our survey, we found in six of the 28 studied systems with an roAp primary one visual companion candidate that has a high probability to be a physical companion, resulting in a multiplicity fraction of $21\\pm9$\\%. This is low compared with similar surveys of A type stars. On the other hand, should all of the objects we believe to be chance projections turn out to be physical companions, the multiplicity fraction may be as high as $46\\pm13$\\%. We introduced a bias in the sample by removing the well studied $\\alpha$\\,Cir, which has a companion. Inserting $\\alpha$\\,Cir in the sample, we get a multiplicity fraction of $24\\pm9$\\%.  Kouwenhoven et al.\\ (\\cite{Kouwenhoven2005}) studied the binarity of A and B stars in the OB association Sco~OB2 with adaptive optics using a $Ks$ filter and a similar field-of-view. 65 of the 199 stars in their sample have at least one companion candidate, leading to a binary fraction of $33\\pm4$\\%. If one restricts the survey to the 113 A type stars, there are 40 stars showing multiplicity, giving a multiplicity fraction of $35\\pm6$\\%. We should note however that the distance to the stars in our sample is on average higher compared with the 130\\,pc to Sco~OB2, with nearly half of the stars in our sample having no parallax determined. Ehrenreich et al.\\ (\\cite{Ehrenreich2010}) find in their volume limited sample of 38 late B to F-type stars, observed with NACO and PUEO, companion candidates to 17 objects. If one restricts their sample to the 19 A-type stars, this leaves companion candidates to seven objects, or a binary fraction of $37\\pm14$\\%. While Ehrenreich et al.\\ employ on some of their observations a larger field-of-view, they consider all companion candidates to A-type stars with separations larger than 7\\arcsec{} as background stars. The distances to these A stars are lower than 67\\,pc, and only two of our own sources would fall into this distance range. Schr\\\"oder \\& Schmitt (\\cite{ SchroederSchmitt2007}) looked at the A-type stars listed in the Bright Star Catalog (Hoff\\/leit \\& Jaschek \\cite{HoffleitJaschek1991}) and found for the 1966 objects listed therein a binary frequency of $16\\pm1$\\%. Looking at a volume limited complete sample of all A stars up to a distance of 50\\,pc, they find 82 binaries among 220 stars, a binary frequency of $37\\pm4$\\%. They identified companion candidates from catalogs, variations in radial velocity or proper motions, variations in ROSAT X-ray light curves at different time scales, and rotational velocity. Since Schr\\\"oder \\& Schmitt did not make use of high spatial resolution observations, they could have missed a number of companion candidates, leading to an underestimation of the true object multiplicity. Assuming that the three different surveys listed above can be summed up, this would lead to 129 binaries for 352 systems, or a binary frequency of $37\\pm3$\\%.  Even taking into account the different observing strategies and object distances used in the studies discussed above, we believe that the 1$\\sigma$ difference hints that our sample shows in comparison a somewhat lower number of stars harboring a companion, assuming that all objects found to have high chance projection probabilities are indeed no physical companions.  Since our sample is quite heterogeneous with respect to distance, we could suffer from a bias that allows us to detect companion candidates only around the nearby roAp stars. We believe that this is not the case. If we separate the sample into three groups with different parallaxes, then we find one companion candidate with low chance projection probability among the five objects with parallaxes above 10\\,mas, two among the ten objects with parallaxes below 10\\,mas, and three among the 13 objects with no parallax measurement, not favoring nearby systems.  In Table~\\ref{tab:multi}, we present the list of the observed roAp stars with notes about their multiplicity. For each object, we indicate whether it is known to be an SB1 and how many astrometric or visual high probability companion candidates are known. Note that there are no known SB2 in our sample. Since the roAp stars have been observed quite extensively for radial velocity variations, the probability that spectroscopic binaries have been missed is quite low. On the other hand, it is not impossible that long term radial velocity variations have been overlooked for individual sources. Of the 28 roAp stars studied, only HD\\,116114 is a potential SB1 system. It is also an astrometric binary. Six objects have visual high probability companion candidates. Another seven stars have visual companion candidates that are very likely chance projections. 15 roAp stars do not show any hint at being part of a binary. We also looked into why the found companion candidates were not detected in radial velocity surveys. For the companion of HD\\,201601, which is the closest we found in our sample with a projected linear separation of 30.1\\,AU, we used the orbital parameters published by Stelzer et al.\\ (\\cite{Stelzer2011}) and computed a radial velocity amplitude of 1.7\\,km\\,s$^{-1}$ for the primary. Kochukhov \\& Ryabchikova (\\cite{KochukhovRyabchikova2001}) found the largest pulsation amplitudes of 0.8\\,km\\,s$^{-1}$ for spectral lines of rare-earth elements. We believe that the accuracy and time span of radial velocity measurements obtained so far is not sufficient to detect radial velocity changes induced by the companion for HD\\,201601. The other likely companions have projected linear separations of 455.1\\,AU (HD\\,99563), 968.4\\,AU (HD\\,101065), and 115.9\\,AU (HD\\,154708AC), which lead to even lower radial velocity amplitudes.  Our recent study (Sch\\\"oller et al.\\ \\cite{Schoeller2010}) of another group of chemically peculiar stars with Hg and Mn overabundances using diffraction-limited near-infrared imaging with NACO led to the detection of 34 near IR companion candidates for the 57 stars studied, confirming that this type of chemically peculiar stars is frequently formed in multiple systems. The interpretation of the difference in duplicity between roAp and HgMn stars and its meaning for the understanding of the origin of the chemical anomalies and pulsations in roAp stars is not straightforward. We assume that most late B-type stars formed in binary systems with certain orbital parameters become HgMn stars (e.g., Hubrig \\& Mathys \\cite{HubrigMathys1995}; Hubrig et al.\\ \\cite{Hubrig2007}; Gonz{\\'a}lez et al.\\ \\cite{Gonzalez2010}; Hubrig et al.\\ \\cite{Hubrig2006}). Our results hint at the possibility that magnetic Ap stars become roAp stars if they are not born in a close binary system.  Tidal forces might conceivably also play a non-negligible role in systems with a larger separation, provided that their eccentricity is large enough. Interaction would then occur mostly on the part of the orbit when the components are closest, since tidal forces are strongly dependent on the distance between the components. At present, though, almost nothing is known about the orbital eccentricities of the noAp binaries. For the roAp star HD\\,201601, Stelzer et al.\\ (\\cite{Stelzer2011}) give an eccentricity of 0.56$\\pm$0.05.   On general grounds, the issue of whether duplicity affects pulsation through tidal interaction is unsettled. From the theoretical point of view, while some authors (e.g., Cowling \\cite{Cowling1941}; Zahn \\cite{Zahn1977}) have conjectured that tides in close binary systems may act as an external perturbing force driving oscillations, the question whether tidal interaction may also be efficient in damping already existing pulsations does not seem to have ever been addressed. Recently, significant advances have been made in tidally-driven pulsations in eccentric binaries with Kepler data. There are 18 heartbeat stars in Thompson et al.\\ (\\cite{Thompson2012}), one of which is KOI-54 (Welsh et al.\\ \\cite{Welsh2011}; Burkart et al.\\ \\cite{Burkart2012}; Fuller \\& Lai \\cite{FullerLai2012}). The theory of Kumar et al.\\ (\\cite{Kumar1995}) fits these heartbeat stars beautifully. Observationally, in the same region of the parameter space in which pulsations were detected, there is only one binary system with a noAp primary presently known, in which the two components are close enough so that significant tidal interaction occurs between them (Giuricin et al.\\ \\cite{Giuricin1984}), the SB1 HD\\,200405 with $P$ = 1.63\\,d (North \\cite{North1998}).  The results of our study support our suspicion that roAp stars are rarely found in binary and multiple systems. However, companionship can not be established based on K photometry alone, and confirming the nature with a near-infrared spectrograph is essential for establishing their true companionship. Future spectroscopic observations in the near-infrared should be used to determine the mass of the companions much more accurately, and explore the physics in their atmospheres by comparison of observed and synthetic spectra.    \\appendix"
    },
    "1208/1208.0341_arXiv.txt": {
        "abstract": "We study the structural evolution of massive galaxies by linking progenitors and descendants at a constant cumulative number density of \\ncum\\ to $z \\sim 3$.  Structural parameters were measured by fitting S\\'{e}rsic profiles to high-resolution CANDELS {\\em HST} WFC3 $J_{125}$ and $H_{160}$ imaging in the UKIDSS-UDS at \\zuds\\ and ACS $I_{814}$ imaging in COSMOS at \\zcosmos.  At a given redshift, we selected the {\\em HST} band that most closely samples a common rest-frame wavelength so as to minimize systematics from color gradients in galaxies.  At fixed $n_{\\rm c}$, galaxies grow in stellar mass by a factor of $\\sim 3$ from $z \\sim 3$ to $z \\sim 0$.  The size evolution is complex: galaxies appear roughly constant in size from $z \\sim 3$ to $z \\sim 2$ and then grow rapidly to lower redshifts.  The evolution in the surface mass density profiles indicates that most of the mass at $r<2$~kpc was in place by $z \\sim 2$, and that most of the new mass growth occurred at larger radii.  This inside-out mass growth is therefore responsible for the larger sizes and higher S\\'{e}rsic indices of the descendants toward low redshift.  At $z<2$, the effective radius evolves with the stellar mass as $r_e \\propto M^{2.0}$, consistent with scenarios that find dissipationless minor mergers to be a key driver of size evolution.  The progenitors at $z \\sim 3$ were likely star forming disks with $r_e \\sim 2$~kpc, based on their low S\\'{e}rsic index of $n \\sim 1$, low median axis ratio of $b/a \\sim 0.52$, and typical location in the star-forming region of the $U-V$ versus $V-J$ diagram.  By $z \\sim 1.5$, many of these star-forming disks disappeared, giving rise to compact quiescent galaxies.  Toward lower redshifts, these galaxies continued to assemble mass at larger radii and became the local ellipticals that dominate the high-mass end of the mass function at the present epoch. ",
        "introduction": "Massive galaxies in the nearby universe are generally comprised of quiescent elliptical and S0 galaxies.  The formation of such massive galaxies has been an active area of study.  Recent observations suggest that the properties of these quiescent galaxies (QGs) were much different at earlier times.  For example, the sizes of QGs have been found to be much smaller, at fixed stellar mass, at high redshift \\citep[e.g.,][]{daddi2005,zirm2007,toft2007,vandokkum2008b,cimatti2008b,vanderwel2008c,franx2008,williams2010,newman2012}.  This implies much higher mass densities within the effective radius for QGs at high redshift.  The size measurements are robust \\citep{szomoru2010,szomoru2012} and the stellar mass measurements are in good agreement with dynamical mass estimates \\citep{cenarro2009,cappellari2009b,vandokkum2009b,vandesande2011}, confirming the dense nature of QGs at high redshift.  Several mechanisms have been proposed to explain the growth in sizes among QGs \\citep[see, e.g.,][]{hopkins2010b}.  Recent discussions have centered on the relative importance of major and minor dissipationless mergers with the latter favored to be the primary channel for size growth \\citep[e.g.,][]{bournaud2007c,naab2009b,bezanson2009,hilz2013}.  While evidence is emerging that these compact QGs represent the cores of local ellipticals \\citep{bezanson2009,hopkins2009j,vandokkum2010}, the progenitors of these compact QGs at even higher redshifts remain a mystery.  Selecting galaxy samples at or above a fixed stellar mass limit has provided important insight into the evolution in properties for such populations.  However, the connection to the evolution of a typical galaxy over cosmic time is not straightforward given that galaxies grow in stellar mass due to in situ star formation and merging: the progenitors of galaxies that lie just above a stellar mass limit at low redshift would not be counted in a census of high redshift galaxies above the same mass limit since they were likely to be less massive.  Measuring the structural evolution of a galaxy as it grows in mass therefore requires a method for linking its progenitors and descendants over cosmic time.  One such method involves selecting galaxies at a constant cumulative number density.  The basic principle behind this method is that the rank ordering of galaxy masses does not change drastically over time.  Therefore, if one selects the 10th most massive galaxy in a comoving volume at $z \\sim 3$, it is still likely to be approximately the 10th most massive galaxy in that comoving volume at $z \\sim 0$, but with a higher overall mass due to star formation and merging.  This technique presents a complementary approach to mass-selected studies.  A number density selection has been used in other recent works to study the structural properties \\citep{vandokkum2010}, star formation histories \\citep[SFHs;][]{papovich2011}, and mass growth \\citep{brammer2011} of galaxies over cosmic time \\citep[see also,][]{loeb2003b}.  The study by \\citet{vandokkum2010} was carried out with ground-based imaging.  In this work, we build on the results in \\citet{vandokkum2010} and utilize high-resolution {\\em Hubble Space Telescope} ({\\em HST}) imaging, allowing us to measure structural properties more accurately and to push to $z \\sim 3$.  At these high redshifts, we also identify the progenitors of $\\sim 2M^{\\star}$ galaxies, which in the local universe are bulge-dominated, quiescent systems with large effective radii.  The layout of this paper is as follows.  In Section~\\ref{sec_data}, we discuss the data used for the study.  In Section~\\ref{sec_analysis}, we discuss the relevant derived quantities.  In Section~\\ref{sec_cumnum}, we examine the structural assembly of a galaxy as it grows in time to become a $\\sim 2 M^{\\star}$ galaxy by $z \\sim 0$.  Our results are further discussed in Section~\\ref{sec_discussion} and we summarize our findings in Section~\\ref{sec_summary}.  For completeness, we briefly overview the structural properties of QGs and star-forming galaxies (SFGs) above a constant stellar mass limit in the Appendix, as this provides an alternative view to the number density selection in the main part of the paper.  We assume a cosmology with $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M = 0.30$, and $\\Omega_{\\Lambda} = 0.70$.  Stellar masses are based on a Chabrier initial mass function \\citep[IMF;][]{chabrier2003}.  Adopting IMFs of different forms below $1$~\\msun\\ \\citep[e.g.,][]{vandokkum2010c} would lead, to first order, to an overall scaling of the stellar masses and otherwise identical results.  All magnitudes are given in the AB system. ",
        "conclusions": "\\label{sec_discussion}    \\subsection{Inside-out Mass Growth}   In Section~\\ref{sec_cumnum}, we have shown that as a galaxy grows in stellar mass by a factor of $\\sim 3$ from $M \\sim 5 \\times 10^{10}$~\\msun\\ at $z=2.75$ to $M \\sim 1.5 \\times 10^{11}$~\\msun\\ at $z \\sim 0$, most of the new stellar mass that is added contributes toward mass growth at larger radii as one moves toward lower redshifts.  The mass profiles presented in Figure~\\ref{fig_cumnum_mass_profile} highlight this point as they show that most of the mass in the core is in place by $z \\sim 2$.  Below $z \\sim 2$, Figure~\\ref{fig_cumnum_mass_rinout} shows how mass growth continues in the outer parts to the present epoch.  As a consequence, the effective radius of the galaxy grows from $r \\sim 2$~kpc at $z \\sim 2-3$ to $r \\sim 7$~kpc at $z \\sim 0$.  Fitting the mass and effective radius evolution at $z<2$, we find that $r_e \\propto M^{2.0 \\pm 0.3}$, which is almost exactly the relation found in \\citet{vandokkum2010} at $z \\lesssim 2$.  The high resolution of the {\\em HST} imaging allowed us to probe the mass distribution of galaxies at $z>2$ at very small radii.  This was not possible in the \\citet{vandokkum2010} study as it relied on ground-based imaging.  With the higher resolution {\\em HST} imaging, we find that the mass within $r=2$~kpc, which is roughly the median effective radius for a galaxy selected at $n_{\\rm c}$ at $z>2$, remains roughly constant at $z<2$.  Thus, there does not appear to be any substantial growth at $z<2$ in the central part of the bulge that will characterize this galaxy at $z \\sim 0$.  From $z \\sim 3$ to $z \\sim 2$, there is evidence for mass growth in the inner part of the galaxy, as seen in Figures~\\ref{fig_cumnum_mass_rinout} and \\ref{fig_cumnum_radii_fixedmass}.  We note, however, that selecting galaxies at different values of $n_{\\rm c}$ can result in a more constant enclosing radius across redshift for a given mass.  Larger samples at $z>2$ may help in clarifying the potential mass buildup in the inner parts of galaxies at those early times.   Dissipationless minor mergers are thought to play a significant role in the buildup of mass for QGs.  In addition, such processes are predicted to assemble mass at large radii, thereby contributing toward the size growth of galaxies \\citep[e.g.,][]{bournaud2007c,naab2009b,bezanson2009}.  For example, using a cosmological hydrodynamical simulation, \\citet{naab2009b} suggest that a $z \\sim 0$ galaxy with $M \\sim 1.5 \\times 10^{11}$~\\msun\\ grew by a factor of $\\sim 3$ after $z \\sim 3$, primarily through dry minor mergers.  This mass growth is almost exactly what is found in this work for the same mass descendant at $z \\sim 0$.  Over this 11~Gyr timespan, the galaxy grew in size by a factor of $\\sim 3-4$, also in agreement with our results.  One caveat in this comparison is that the overall sizes of the simulated galaxy are a factor of $2-3$ lower than what is found in this work.  With more recent simulations, \\citet{hilz2013} show that the relation $r_e \\propto M^{2}$ \\citep[see also,][]{laporte2012}, found in this work and in \\citet{vandokkum2010}, is most easily reproduced by $\\sim 3-5$ mergers with mass ratios of 1:5.  In contrast, major mergers lead to a close to linear dependence of $r_e$ on $M$, which is not supported by the observations.  It remains to be seen whether these mergers are actually observed \\citep[see, e.g.,][]{williams2011,newman2012}.   \\subsection{The Progenitors of Local Ellipticals at $z \\sim 3$}   In the local universe, the most massive systems are generally bulge dominated elliptical galaxies \\citep[e.g.,][]{vanderwel2009b,holden2012} with large effective radii \\citep[e.g.,][]{shen2003}.  In this work and in \\citet{vandokkum2010}, we showed that the progenitors of these local ellipticals at $1 \\lesssim z \\lesssim 2$, are primarily compact QGs.  As discussed in the previous section, these compact QGs are generally considered the cores of ellipticals at $z \\sim 0$ \\citep[see also,][]{hopkins2009j}, growing in mass and size potentially through dissipationless minor mergers to match the properties, namely sizes, of local ellipticals.  With the {\\em HST} imaging and the deep-IR UDS data, our number density selection allows us to the trace the properties of progenitors of local ellipticals to $z \\sim 3$.  At these redshifts, Figures~\\ref{fig_uvj_ncum} and \\ref{fig_cumnum_qg_frac} show that most of the progenitors are SFGs.  The S\\'{e}rsic indices in Figure~\\ref{fig_cumnum_structure}(b) suggest that these SFGs have exponential profiles, which is typically associated with the surface brightness profiles of disks.  The S\\'{e}rsic index alone, however, is not definitive in defining the shapes of these progenitors at $z \\sim 3$.  Instead, the most compelling evidence comes from the axis ratios shown in Figure~\\ref{fig_cumnum_structure}(c).  At $z \\sim 0$, galaxies selected at \\ncum\\ have median observed axis ratios of $b/a \\sim 0.75$.  The implied intrinsic axis ratio is roughly 2:3 \\citep{holden2012}, indicative of spheroidal systems.  Meanwhile, at $z \\sim 3$, Figure~\\ref{fig_cumnum_structure}(c) shows that galaxies at $n_{\\rm c}$ have a median axis ratio of $b/a \\sim 0.52$.  This value for the axis ratio is very low considering that for a population of randomly oriented, infinitely thin disks, the median axis ratio would be $b/a \\sim 0.5$.  It is therefore likely that the progenitors of massive, quiescent, local elliptical galaxies at $z \\sim 3$ are star-forming disks.  Many of the star-forming disks at $z \\sim 3$ disappear over the redshift range $1.5<z<3$ and give way to compact QGs, as seen in Figure~\\ref{fig_uvj_ncum}.  While the details of this transition require further investigation, we note that for a much broader mass range than what is considered here, \\citet{barro2012} identify a subset of the star forming population at $2<z<3$ with similar structure to that of compact QGs \\citep[see also,][]{stefanon2013}.       We used {\\em HST} imaging to study the structural properties of galaxies selected at a constant cumulative number density of \\ncum\\ at redshifts of \\zwindow.  This selection allowed us to trace the evolution of galaxies with stellar mass $M=5 \\times 10^{10}$~\\msun\\ at $z=2.75$, as they grew by a factor of $\\sim 3$ to become $\\sim 2M^{\\star}$ galaxies in the local universe.  This work builds on the previous analysis by \\citet{vandokkum2010}, who also selected galaxies at a constant number density.  Here, we employ high-resolution {\\em HST} imaging and extend the analysis to $z \\sim 3$.  In contrast to mass-selected studies, our selection at a constant cumulative number density allows for a more straightforward evolutionary link between progenitors and descendants at different redshifts.  At \\zuds, we employed CANDELS WFC3 $J_{125}$ and $H_{160}$ imaging in the UDS, while at \\zcosmos\\ we used wide-field ACS $I_{814}$ imaging in COSMOS to fit single-component S\\'{e}rsic profiles at a common rest-frame wavelength.  The resulting S\\'{e}rsic indices, effective radii, and axis ratios were used to aid in our analysis.  The uniform data sets and analysis methods carried out in this work serve to minimize systematics.  Our main conclusions are the following:  \\begin{enumerate}  \\item The typical galaxy at the selected value of $n_{\\rm c}$ has grown in effective radius by a factor of $\\sim 3-4$, mostly since $z \\sim 2$ (Figure~\\ref{fig_cumnum_structure}).  \\item The evolution in the stellar mass surface density profiles of galaxies selected at $n_{\\rm c}$ indicates that most of the mass in the central regions was in place by $z \\sim 2$, while almost all of the new mass growth took place in the outer parts (Figures~\\ref{fig_cumnum_mass_profile}-\\ref{fig_cumnum_radii_fixedmass}).  This inside-out mass growth is responsible for the increase in size and S\\'{e}rsic index toward low redshift.  \\item At $z<2$, we find that as the stellar mass builds up, the effective radius grows as $r_e \\propto M^{2.0}$, in excellent agreement with \\citet{vandokkum2010}.  Recent simulations show that such a dependence is consistent with mergers with mass ratios of 1:5 being responsible for much of the size growth \\citep{hilz2013}.   \\item At $z \\sim 3$, the rest-frame $UVJ$ colors, S\\'{e}rsic indices, and axis ratios indicate that the progenitors of present day massive galaxies were star-forming disks with $r_e \\approx 2$~kpc and a third of the $z \\sim 0$ stellar mass.  At $1.5 \\lesssim z \\lesssim 2$, these galaxies doubled in stellar mass and were mostly compact QGs.  These galaxies evolved further into the $\\sim 2M^{\\star}$ galaxies in the local universe that are known to be quiescent, bulge dominated, elliptical galaxies with large effective radii.   \\end{enumerate}"
    },
    "1208/1208.1467.txt": {
        "abstract": "{Theoretical spectra of X-ray bursting neutron star (NS) model atmospheres are widely used to determine the basic NS parameters such as their masses and radii. Compton scattering, which plays an important role in spectra formation at  high luminosities, is often accounted for using  the differential Kompaneets operator, while in other models a more general,  integral operator for the Compton scattering kernel is used. } {We  construct  accurate NS atmosphere models using for the first time an exact treatment of Compton scattering via the integral relativistic  kinetic equation. We also  test various approximations to the Compton scattering redistribution function and compare the results with the previous calculations based on the Kompaneets operator.  } {We solve the radiation transfer equation together with the hydrostatic equilibrium equation accounting exactly for the radiation pressure by electron scattering. We use the exact relativistic  angle-dependent redistribution function as well as its simple approximate representations. } {We thus construct a new set of plane-parallel atmosphere models in local thermodynamic equilibrium (LTE) for hot NSs. The models were computed for six chemical compositions (pure H, pure He, solar H/He mix with various heavy elements abundances $Z$ = 1, 0.3, 0.1, and 0.01$Z_{\\odot}$, and three surface gravities $\\log g$ = 14.0, 14.3, and 14.6. For each chemical composition and surface gravity, we compute more than 26 model atmospheres with various  luminosities relative to the Eddington luminosity $\\leddth$ computed for the Thomson cross-section. The maximum relative luminosities $L/\\leddth$ reach values of up to 1.1 for  high gravity models. The emergent spectra of all models are redshifted and fitted by diluted blackbody spectra in the 3--20 keV energy range appropriate for the {\\it RXTE}/PCA. We also compute the color correction factors $\\fc$. %\\footnote{Tables D.1, D.2 and D3 %containing spectra and color correction factors are only available in electronic form at the CDS %via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) %or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/.} } {The radiative acceleration $g_{\\rm rad}$ in our luminous, hot-atmosphere models is significantly smaller than in corresponding models based on the Kompaneets operator, because of the Klein-Nishina reduction of the electron scattering cross-section, and therefore formally ``super-Eddington'' model atmospheres do exist. The differences between the new and old model atmospheres are small for $L/\\leddth< 0.8$. For the same $g_{\\rm rad}/g$, the new $\\fc$ are slightly larger (by approximately 1\\%)  than the old values. We also find that the model atmospheres,  the emergent spectra, and the color correction factor computed using angle-averaged and approximate Compton scattering kernels differ from the exact solutions by less than 2\\%. } ",
        "introduction": "X-ray bursting neutron stars (NSs) are members of low-mass X-ray binaries with quasi-periodical thermonuclear flashes on their surfaces \\citep[see reviews by][]{LvPT93, SB06}. Thermonuclear burning occurs  at the bottom of the freshly accreted matter and can be so powerful that the luminosity reaches the Eddington limit. These bursts lead to photospheric radius expansion (PRE) and are a potentially powerful tool for determining  the NS masses and radii \\citep{Ebisuzaki87,Damen:90,vP:90}. The knowledge of NS basic parameters is extremely important for establishing the  physical properties (equation of state) of supra-nuclear  dense matter in the NS inner cores  \\citep[see][for a review]{LP07}.  The precise radius determination of the NSs from X-ray bursts is  impossible without accurate spectral models. The observed spectra of X-ray bursts are often well-fitted by a blackbody  \\citep{gallow08}. Theoretical models of hot NS atmospheres show  \\citep{Londonetal:86,Lapidus:86,Pavlov.etal:91} that the emergent spectra are close to diluted blackbody spectra $F_{E} \\approx \\fc^{-4} B_{E} (\\fc T_{\\rm eff})$ owing to strong energy exchange caused by Compton scattering. The color correction factor $\\fc \\equiv T_{\\rm c}/ T_{\\rm eff}$, which relates the color temperature of the spectrum $T_{\\rm c}$ to the  effective temperature of the atmosphere $T_{\\rm eff}$, takes values of about 1.3--1.9. Theoretical dependences of the color correction factor on luminosity for various chemical compositions and gravities were computed by  \\citet[][hereafter SPW11]{SPW11}. The atmosphere models used in their work were based  on the Kompaneets operator \\citep{Kompaneets:57}  description of Compton scattering. This approach was also used in most previous studies \\citep{Londonetal:86,Lapidus:86,Ebisuzaki87,Pavlov.etal:91}. The Kompaneets operator describes Compton scattering in the non-relativistic, isotropic, diffusion approximation, which should still be rather adequate for hot NS model atmospheres with typical effective temperatures below $\\sim$2 keV. On the other hand, Madej and collaborators \\citep{madej:91,Madej.etal:04,Majczyna:05} used the integral description of Compton scattering employing the angle-averaged redistribution function (RF) derived by \\citet{Guilbert81}. A fully relativistic treatment of Compton scattering becomes important at luminosities close to the Eddington values, because of the reduction of the effective cross-section and the corresponding drop in the radiation pressure force. This obviously changes the maximum value of the luminosity where hydrostatic equilibrium can still be achieved, i.e. the actual value of the Eddington luminosity. This has important implications for the method of determining the NS masses and radii from PRE bursts,  which is based on the Eddington luminosity (see \\citealt{S11}, hereafter S11). It is therefore,  necessary to check the validity of model atmospheres and spectral properties computed with the Kompaneets operator by comparing the results with more accurate models based on an exact treatment of Compton scattering  using relativistic kinetic equations.  Here we construct new models of hot NS atmospheres based on an exact treatment of Compton scattering that employs a fully relativistic, angle-dependent RF  \\citep{AA81,PKB86,NP94,PS96}. In Sect. \\ref{s:methods}, we  present the main equations describing the atmosphere models as well as the methods for their solution. We also compare  atmosphere models based on various RFs for Compton scattering and using the Kompaneets operator. In Sect. \\ref{s:newgrid}, we present a new set of atmosphere models, emergent spectra, and the color correction factors. We compare them to previous results obtained with the Kompaneets operator. In Sect.~\\ref{s:application}, we discuss applications of the new models to the observations of X-ray bursts and determination of the NS masses and radii. We summarize in Sect.~\\ref{sec:summary}. In Appendix~\\ref{sec:rke} we derive the exact RF for Compton scattering. The details of the method of solution of the radiative transfer equation are presented in Appendix~\\ref{sec:method_rte}. Comparison with the previous attempts to compute atmosphere models using integral approach to Compton scattering are given in Appendix~\\ref{sec:madej}. And, finally, Appendix~\\ref{app:D} presents the spectral characteristics of the models. ",
        "conclusions": "\\label{sec:summary}  We  have considered hot NS model atmospheres taking into account Compton scattering using the relativistic kinetic equation with the exact angle-dependent RF and  cross-section \\citep{NP94,PS96}, accounting also for the induced scattering. We have developed a method to solve the obtained radiation transfer equation using a short characteristic method with the accelerated $\\Lambda$-iteration. We have implemented this solution in our computer code for the model atmosphere calculations (SPW11).  We have examined the properties of the  new model atmospheres. The main difference in comparison with the old models computed with  the Kompaneets operator concerns radiative acceleration. In the new approach, the Klein-Nishina reduction in the electron scattering cross-section leads to a decrease in the radiative acceleration  relative to that for  the Thomson  scattering cross-section. The grid of model atmospheres was extended to higher effective temperatures up to  luminosities formally exceeding the Eddington limit $\\leddth$ computed for the Thomson cross-section. %The radiation pressure decreases that occur in the deeper, hot, optically thick layers are due to the same effect. The role of the radiation pressure reduces in the deeper, hot, optically thick layers due to the same effect.  We computed a new set of 484 hot NS model atmospheres. Following SPW11, we computed the models for six different chemical compositions (pure hydrogen, pure helium, solar hydrogen/helium mix with various abundances of heavy elements $Z$ = 1, 0.3, 0.1, and 0.01 $Z_{\\odot}$) and three surface gravities ($\\log g =14.0, 14.3$, and 14.6). The relative luminosities range from $l$=0.001 to 1.06--1.1 (depending on $\\log g$). The new models with $l < 0.8$ are almost identical to the old models based on the Kompaneets operator approach. At higher $l$, the deviations are more significant due to a different treatment of the radiative acceleration. The difference between the new and old models is small if instead of $l$ one uses the relative radiative acceleration $g_{\\rm rad}/g$ as a parameter.  The spectra of the new and old models deviate by less than 5\\% in the region around the spectral peak, and the deviation grows at higher energies. We have also tested various approximate and angle-averaged RFs for Compton scattering and found that they produce spectra that deviate from the exact solution typically by less than 2\\%. The comparison of our spectra to the models computed using Madej's code \\citep{madej:91,Madej.etal:04,Majczyna:05}  revealed dramatic differences in the spectral shape. We attribute these differences to their usage of the incorrect RF from \\citet{Guilbert81} and/or to the non-convergence of their radiation transfer calculations.  The spectra of all our models were fitted by the diluted blackbody spectra in the {\\it RXTE}/PCA energy band (3--20 keV) using the same five fitting procedures as in SPW11. The new $\\fc$--$l$ relations are similar in shape to the old relations, but depend significantly on the surface gravity. However, the dependences of the color correction on $g_{\\rm rad}/g$ are almost indistinguishable, with the new $\\fc$ being larger by  approximately 1\\%. The color corrections with the corresponding dilution factors, the theoretical emergent spectral energy distributions, and specific intensities at various angles for all models are available at CDS.  In the present paper, we have also tested how the new  $\\fc$--$l$ relations affect the determination of the NS masses and radii derived from the  spectral evolution observed during the cooling stages of PRE bursts. We found that the best-fit NS radius increases by about 10\\% from the value obtained using the old model based on the Kompaneets equation (S11). We note that bursting NSs are rapidly rotating and accounting for the latitudinal variation in gravity and the Doppler effect will affect the estimated NS radius.  This will be a subject of a separate publication."
    },
    "1208/1208.1347_arXiv.txt": {
        "abstract": "The gamma-ray binary system \\psrbl\\ consists of a 48~ms pulsar orbiting a Be star. The system is particularly interesting because it is the only gamma-ray binary system where the nature of the compact object is known.  The non-thermal radiation from the system is powered by the spin-down luminosity of the pulsar and the unpulsed radiation originates from the stand-off shock front which forms between the pulsar and stellar wind. The Be star/optical companion in the system produces an excess infrared flux from the associated circumstellar disc. This infrared excess provides an additional photon source for inverse Compton scattering. We discuss the effects of the IR excess near periastron, for anisotropic inverse Compton scattering and associated gamma-ray production.  We determine the infrared excess from the circumstellar disc using a modified version of a curve of growth method, which takes into account the changing optical depth through the circumstellar disc during the orbit. The model is constrained using archive data and additional  mid-IR observations obtained with the VLT during January 2011. The inverse Compton scattering rate was calculated for three orientations of the circumstellar disc. The predicted gamma-ray light curves show that the disc contribution is a maximum around periastron and not around the disc crossing epoch. This is a result of the disc being brightest near the stellar surface.  Additional spectroscopic and near-infrared observations were obtained of the system and these are discussed in relation to the possibility of shock heating during the disc crossing epoch. ",
        "introduction": "The gamma-ray binary system \\psrbl\\ consists of a 48~ms pulsar orbiting around a Be star.  The orbit is eccentric  ($ e \\approx 0.87 $) with an orbital period of  $P_\\rmn{orb}\\sim3.4$~yr \\citep{johnston92,johnston94}.  The system is particularly interesting because it is the only gamma-ray binary system where the nature of the compact object is known.  Unpulsed TeV gamma-ray emission has previously been detected around periastron by the  High Energy Stereoscopic System (H.E.S.S.) \\citep{aharonian05,aharonian09}.  Previous observations have also detected unpulsed radio emission near periastron \\citep[e.g][]{johnston99} and  X-ray emission across the whole orbit \\cite[see e.g.][]{cominsky94full,kaspi95,chernyakova06,chernyakova09}.  The discovery of extended radio and X-ray structures have been recently reported \\citep*{moldon11, pavlov11}.  These observations have shown that the unpulsed emission is variable at all wavelengths with respect to the orbital period. The system was observed around the 2010 periastron passage ($\\sim$ 14 December 2010) with  {\\em Fermi}/LAT in the high energy (HE, $0.1-100$~GeV) regime and showed a flux modulation during the periastron passage, peaking $\\sim 30$~days after periastron \\citep{abdo11, tam11}.  The non-thermal radiation from the system is powered by the spin-down luminosity of the pulsar and the unpulsed radiation originates from the stand-off shock front which forms between the pulsar and stellar wind.  Extensive modelling of this process has been presented in \\citet{tavani97}, where the X-rays are produced via synchrotron radiation of ultra-relativistic electrons (Lorentz factor $\\gamma \\sim 10^6$) and gamma-rays through the inverse Compton (IC) scattering of target photons from the Be star.  The ratio between the X-ray and gamma-ray flux is dependent on the radiative and adiabatic cooling times of the post-shocked pulsar wind. IC cooling was predicted to increase near periastron.  IC cooling of the pre-shocked pulsar wind was also suggested for \\psrb\\ by \\citet{ball00}.  The optical companion in the system, LS~2883,  is a Be star, and produces an infrared (IR) flux which is higher than predicted by Kurucz stellar atmosphere models \\citep{kurucz79}.   This IR excess is produced by free-free and free-bound emission in the circumstellar disc which surrounds the Be star.  These Be star circumstellar discs are decretion discs which form around Be stars and are regions of enhanced stellar outflow.  Optical interferometry observations of these extended circumstellar envelopes have shown that they are symmetrical around the rotation axis \\citep[see e.g.][]{quirrenbach94}. The decretion discs are variable and grow and shrink over periods of hundreds to thousands of days.  See e.g.\\ \\citet{porter03} for a review of classical Be stars.  The IR excess from the circumstellar disc provides an additional photon source to the blackbody flux from the optical star.  \\citet{vansoelen11} showed that in the isotropic case this IR excess could increase the gamma-ray flux, particularly at GeV energies, by a factor $\\sim 2$.   It is possible that near the disc crossing the IR excess may have a much greater effect. In this study we discuss the effects of the IR excess near periastron, adopting an anisotropic inverse Compton model \\citep*[e.g.][]{fargion97,dubus08}.  The soft photon excess provides target photons for IC scattering in the Thomson limit where the scattering cross section is a maximum. The additional effects of the different electron cooling processes in the plerion shock front \\citep*[e.g.][]{kirk99}, as well as the effects of gamma-ray absorption \\citep{dubus06}, will not be discussed in this paper and the reader is directed to the given reference for relevant discussions. In order to include the IR excess, we determine the target photon distribution using a modified version of the curve of growth method which was first proposed by \\citet{waters86} (referred to hereafter as the COG method).  This modified COG model takes into account the changing optical depth as seen by the pulsar during its orbit.  We constrain the COG model using literature data and VLT observations which were taken during January 2011, approximately 22~days after the December 2010 periastron passage.  This paper is structured as follows: an outline of previous observations and modelling of \\psrb\\ is presented in section~\\ref{psrb}; section~\\ref{sec:observations} discusses observations we obtained with the VLT, SALT,  SAAO~1.9m~ and IRSF~telescopes after the 2010 periastron passage; section~\\ref{be-stars} discusses the modified COG method that we have derived;  the anisotropic IC scattering is discussed in section~\\ref{ic}; section~\\ref{sec:modelling} outlines our modelling of the anisotropic inverse Compton scattering around periastron; sections~\\ref{sec:discussion} and \\ref{sec:conclusion} discuss the results and implications thereof. ",
        "conclusions": "\\label{sec:conclusion}   The binary star system \\psrbl\\ has remained a fascinating source since the detection of the pulsar approximately two decades ago.  The interaction between the Be star and the pulsar produces multi-wavelength emission from radio up to TeV gamma-rays, covering 19 decades in energy.   The optical companion in the system is a Be type star, known to possess a circumstellar disc which produces an IR flux.  This flux is observed as an IR excess which is higher than predicted by blackbody and stellar atmosphere models.  Since the gamma-rays in the system are produced by the up-scattering of photons, the presence of an extended and additional photon source could potentially have a significant effect on the gamma-ray production in \\psrbl.  We modified the COG method, first proposed by \\citet{waters84}, which allowed us to predict the flux originating from the circumstellar disc at any position near the disc.  This method has been used to predict the photon density as observed from the point-of-view of the pulsar.  The COG method was constrained using optical and IR data obtained from literature and VLT observations.  The photon density was calculated around periastron and we applied it to the problem of AIC scattering.  The anisotropic scattering considered the change in the photon distribution during the orbit, the variation in the solid angle subtending the disc, and the direction of scattering.  The variation in the gamma-ray spectrum around periastron was calculated for different disc orientations.  In this study the leptons in the pulsar wind are assumed to be mono-energetic before the stand-off shock and a power-law energy distribution in the post-shock region.  Both a pre-- and post--shock distribution have been considered. For \\psrb, it can be shown that adiabatic cooling dominates near the stand-off shock and, for this reason, only adiabatic cooling was considered for the AIC modelling.  The IC modelling showed that the inclusion of the IR excess associated with the circumstellar disc increased the GeV gamma-rays by a factor $\\ga 2$.  This relative increase in the flux is partly dependent on the choice of the electron distribution. For a wide distribution (e.g.\\ $10^4 \\leq \\gamma \\leq 10^7$), as was shown for the isotropic modelling \\citep[fig.~3 in][]{vansoelen11}, or an electron distribution centred around $\\gamma \\sim 10^4$ (Fig.~\\ref{fig:fermi_multiplot_preshock}), the IC scattering is dominated by the scattering of stellar photons in the Thomson limit.  However, this study has shown that for the electron distribution proposed by \\cite{kirk99}, which has been shown to provide a reasonable fit to the 2004 H.E.S.S.\\ observations, the IR excess from the circumstellar disc increases the gamma-ray flux by a factor $\\geq 2$ at GeV energies.   The predicted light curves show (Figs.~\\ref{fig:aic_light_curve_fermi_new}, \\ref{fig:aic_light_curve_hess_new} \\& \\ref{fig:fermi_multiplot_preshock}) that the circumstellar disc is most influential close to periastron and, generally, after the disc crossing. This is a result of the disc being brightest near the stellar surface  and the associated favourable scattering angles after the disc crossing.  While it was thought the predicted gamma-ray flux may be highest near the disc crossing epoch ($\\sim \\tau \\pm 20$~days), the low intensity of the disc at $\\sim 50$ stellar radii from the star's surface does not provide a sufficient number of photons to increase significantly the gamma-ray flux in this region.  This study has only considered a steady state disc and has not considered what effects the pulsar passage may have introduced on the circumstellar disc.  Shock heating of the disc or clumping of material may result in regions of higher intensity which could result in more localised flux increases.  The additional effects such as changes in the electron cooling have not been considered, as these have been addressed by other authors.  This study has focussed on the effect of the IR excess from the circumstellar disc in order to determine its influence on the production of GeV energy gamma-rays.  The study shows that the gamma-ray flux could be increased by a factor greater than two at GeV energies, which is a significant increase.  However, the predicted gamma-ray flux lies far below that detected by {\\em Fermi}/LAT  at the previous periastron passage. A combination of a number of effects, including e.g.\\ Doppler boosting, shock heating and changes in the shock front are required to explain the flare that was observed by {\\em Fermi}/LAT  approximately one month after periastron."
    },
    "1208/1208.5943_arXiv.txt": {
        "abstract": "Low-mass objects embedded in isothermal protoplanetary discs are known to suffer rapid inward Type I migration.   In non-isothermal discs, recent work has shown that a  decreasing radial profile of the disc entropy can lead to a strong positive corotation torque which can significantly slow down or reverse Type I migration in laminar viscous disc models, depending on the amount of viscous and thermal diffusion operating  in the planet's horseshoe region. Since the latter is a fraction of the pressure scale height of the disc, it is not clear however how this picture changes in turbulent disc models. The aim of this study is to examine the impact of turbulence on the torque experienced by a low-mass planet embedded in a non-isothermal protoplanetary disc. We particularly focus on the role of turbulence on the corotation torque whose amplitude depends on the efficiency of diffusion processes in the planet's horseshoe region. The main issues we want to address are whether the part of the corotation torque scaling with the entropy gradient can remain unsaturated in the presence of turbulence and whether the saturation process in non-isothermal discs can be satisfactorily modelled using laminar disc models. We performed 2D numerical simulations using a grid-based hydrodynamical code in which turbulence is modelled as stochastic forcing. In order to provide estimations for the viscous and thermal diffusion coefficients as a function of the amplitude of turbulence,  we first set up non-isothermal disc models for different values of the amplitude of the turbulent forcing. We then include a low-mass planet and determine the evolution of its running time-averaged torque. We show that in non-isothermal discs, the entropy-related corotation torque can indeed remain unsaturated in the presence of turbulence.  For turbulence amplitudes that do not strongly affect the disc temperature profile, we find that the running time-averaged torque experienced  by an embedded protoplanet is in  fairly good agreement with laminar disc models with appropriate values for the thermal and viscous diffusion coefficients and with the formulae of Paardekooper et al. (2011) for the total torque in non-isothermal discs. In discs with turbulence driven by stochastic forcing, the corotation torque therefore behaves  similarly as in laminar viscous discs and can be responsible for significantly slowing down or reversing Type I migration. ",
        "introduction": "In the course of its evolution, the tidal interaction of a planet with the protoplanetary disc in which it formed is known to cause a significant change of the planet orbital elements.  In particular, low-mass planets can experience dramatic changes of their distance from the central star under the process of Type I migration (Ward 1997; Tanaka et al. 2002). Type II migration concerns more massive planets that are able to strongly modify the underlying disc structure by creating a gap around their orbit (Lin \\& Papaloizou 1986). In that case, a planet undergoing Type II migration is locked to the disc viscous evolution and migrates inward on a timescale corresponding to the disc viscous timescale, provided that the characteristic disc mass is still larger than the planet mass (Nelson et al. 2000).\\\\ The gravitational torque exerted by the disc on the planet and causing Type I migration  consists of two components. The differential Lindblad torque results from the angular momentum exchange between the planet and the spiral density waves it generates inside the disc.  For sufficiently low-mass planets, the Lindblad torque is well described by a linear analysis, which shows that the specific value of this torque scales linearly with disc mass and planet mass and as the inverse square of the disc aspect ratio (Tanaka et al. 2002). Although its sign depends on the  density and temperature gradients inside the  disc, the differential Lindblad torque is generally negative for typical disc models and is therefore responsible for inward migration.\\\\ The corotation torque is due to the torque exerted by the material located in the coorbital region of the planet. A linear analysis reveals that the corotation torque consists of a barotropic part  which scales with the vortensity (i.e. the ratio between the vertical component of the vorticity and the disc surface density) gradient (Goldreich \\& Tremaine 1979) plus an entropy-related part which scales with the entropy gradient (Baruteau \\& Masset 2008; Paardekooper \\& Papaloizou 2008). A negative vortensity (resp. entropy) gradient gives rise to a positive vortensity (resp. entropy) related corotation torque. The consequence is that for midly positive surface density gradients or negative entropy gradients, a positive  corotation torque can eventually counteract the effect of a negative differential Lindblad torque,  which may stall or even reverse migration (Masset et al. 2006; Paardekooper \\& Papaloizou 2009). In isothermal discs, the corotation torque is a non-linear process generally referred as the horsehoe drag and whose amplitude is controlled by advection-diffusion of vortensity inside the horsehoe region. In non-isothermal discs, singular production of vortensity also arises due to the entropy discontinuity on downstream separatrices (Masset \\& Casoli 2009; Paardekooper et al. 2010) and in that case, the amplitude of the corotation torque is therefore powered by advection-diffusion-production of vortensity inside the horseshoe region.\\\\ In the absence of any diffusion processes inside the disc, vortensity and entropy gradients across the horseshoe region tend to flatten through phase mixing, which causes the two components of the horseshoe drag to  saturate. Consequently, desaturating the horseshoe drag requires that some amount of viscous and thermal diffusions are operating inside the horseshoe region. In that case, the amplitude of the horseshoe drag depends on the ratio between the diffusion timescales and the horseshoe libration timescale and its optimal value, also referred as the fully unsaturated horseshoe drag, is obtained when the diffusion timescales are approximately equal to half the horseshoe libration time (see e.g. Baruteau \\& Masset 2012 for a recent review). In the limit where the diffusion timescales become shorter than the U-turn timescale, the corotation torque decreases and approaches the value predicted by linear theory. Therefore, the corotation torque can be considered as a linear combination of the fully unsaturated horseshoe drag and the linear corotation torque with coefficients depending on the ratio between the diffusion timescales and the horseshoe libration timescale. Corotation torque formulae as a function of viscosity and thermal diffusivity were recently proposed by Paardekooper et al. (2011) and Masset \\& Casoli (2010). \\\\ So far, most of the studies aimed at studying the process of saturation of the corotation torque focused on laminar disc models with prescribed values of kinematic viscosity and thermal diffusivity. However, it is likely that both angular momentum and heat transport in protoplanetary discs occur because of turbulence probably due to the magneto-rotational instability (MRI, Balbus \\& Hawley 1991) and it is not clear how the disc turbulence really impacts the corotation torque. In the case of isothermal discs, Uribe et al. (2011) performed 3D MHD simulations of planet migration in stratified discs that suggested the torques coming from the corotation region can remain unsaturated due to viscous stresses in the disc. Baruteau et al. (2011) recently confirmed the existence of an unsaturated corotation torque and demonstrated that a horseshoe dynamics is indeed at work in turbulent discs with a weak toroidal magnetic field. Using 2D hydrodynamical simulations in which disc turbulence is modelled through stochastic forcing, Baruteau \\& Lin (2010) showed that in isothermal discs, the running time-averaged torque experienced by a protoplanet is in good agreement with that obtained from a laminar disc model with a similar vortensity diffusion coefficient, and  in which the total torque consists of  the Lindblad torque plus the  barotropic part of the corotation torque. \\\\ In this paper we adopt a similar strategy as in Baruteau \\& Lin (2010) but focus on non-isothermal discs in which the horseshoe drag has an additional contribution due to the presence of an entropy gradient across the horseshoe region. The aim is to investigate whether the entropy-related horseshoe drag remains unsaturated in presence of turbulence and whether the running time-averaged torques in turbulent runs are in good agrement with a laminar viscous disc model with appropriate values for the vortensity and entropy diffusion coefficients. \\\\ In the context of planetary population synthesis models aimed at reproducing the  mass-semi-major axis diagram of exoplanets, it is of crucial importance to investigate whether in turbulent disc models the entropy-related corotation torque behaves as in laminar viscous discs since it may be responsible for significantly slowing down or reversing Type I migration in the latter case. \\\\ The paper is organized as follows. In Sect. 2, we give the basic equations and notations we use in this work and describe the numerical set-up in Sect. 3. In Sect. 4, we present the properties of turbulence in our simulations. In particular, we evaluate the vortensity and entropy diffusion coefficients as well as the turbulent Prandtl number. In Sect. 5, we estimate the effects of turbulence on the entropy-related corotation torque and compare with the torque formulae of Paardekooper et al. (2011) and with laminar calculations as well. We finally discuss our results and draw our conclusions in Sect.6. ",
        "conclusions": "In this paper, we have presented the results of 2D hydrodynamical simulations which study the interaction of an embedded low-mass planet with a turbulent non-isothermal disc. The aim is to investigate the role of turbulence on Type I migration in non-isothermal discs, with particular emphasis on the process of desaturation of the entropy-related corotation torque by turbulence. \\\\ We employ the turbulence model of Laughlin et al. (2004) in which a turbulent potential corresponding to the superposition of multiple wave-like modes is applied to the disc. This turbulence potential  had been shown previously to reproduce the main statistical properties of MHD turbulence in isothermal discs (Baruteau \\& Lin 2010) . In non-isothermal discs, turbulent density and entropy fluctuations generate transport of momentum and heat inside the disc and we show that both the associated viscous and thermal diffusion coefficients scale as $\\gamma^2$, where $\\gamma$ refers to the amplitude of the turbulent forcing. This results in a turbulent Prandtl number $P_r=\\alpha_R c_s H/\\kappa$ where $\\alpha_R$ is the Reynolds stress parameter  which is almost constant with $\\gamma$ and we estimate for the first time that $P_r \\sim 0.3$ in non-isothermal turbulent discs.  We caution, however, that such a value for the turbulent Prandtl number should be checked against 3D MHD, non-isothermal simulations in which turbulence is generated by the MRI. \\\\ By investigating the evolution of the running time-averaged torque experienced by a protoplanet for different values of $\\gamma$, we show that turbulence can unsaturate the entropy-related corotation torque.  For similar values for the viscous and thermal diffusion coefficients, we find that the formulae of Paardekooper et al. (2011) for non-isothermal Type I migration in laminar disc models reproduce quite satisfactorily the time-averaged torques deduced from our turbulent runs,  provided that the typical amplitude of turbulent fluctuations is not too large so that turbulent transport of heat and momentum do not significantly modify the background density and temperature profiles. For high values of $\\gamma$  corresponding to  values for the dimensionless $\\alpha$ viscosity of Shakura \\& Sunyaev (1973) $\\alpha \\gtrsim 10^{-3}$, we indeed observe discrepancies between the analytical estimate and the results of simulations which, to a large extent, are related to the alteration of the temperature profile by turbulence. We find also that the amplitudes of the running time-averaged turbulent torques  are slightly smaller than the torques amplitudes deduced from laminar simulations with initial density and temperature profiles corresponding to the time averaged profiles resulting from turbulent runs. This suggests that the entropy-related horseshoe drag in presence of turbulence is slightly more saturated compared to an equivalent laminar disc model with similar vortensity and entropy diffusion coefficients. \\\\ Although these results were obtained from a disc model in which the initial entropy gradient is positive, similar conclusions  should be drawn in the case of a negative initial entropy gradient. Since in that case the amplitude of the total torque is only a small fraction of the typical amplitude of the stochastic torque, turbulent runs would require a much longer convergence time than the simulations presented here. For a negative initial entropy gradient, the entropy-related corotation torque is positive and can eventually counteract the effect of the (negative) differential Lindblad torque, depending on the local temperature gradient and provided that the viscous and thermal diffusion coefficients are large enough to keep the entropy-related corotation torque from saturating. \\\\ The positive results obtained in this paper bring further support for the entropy-related corotation torque as a possible solution for the discrepancy between the mass-semi-major axis diagram of observed exoplanets, and that inferred from theories of planetary formation and evolution (e.g., Ida \\& Lin 2008, Mordasini et al. 2009, Howard et al. 2010, Hellary \\& Nelson 2012).\\\\ The simulations we have presented have neglected radiative processes and we will present in a future paper simulations  with radiative transfert treated in the flux-limited approximation. We will also investigate the effects of turbulence on the dynamics of multiple protoplanets located in the vicinity of an opacity transition where an entropy jump can act as a planet trap and promotes embryo growth (Horn et al. 2012)."
    },
    "1208/1208.2939_arXiv.txt": {
        "abstract": "{We present a toy model of mildly super-Eddington, optically thin accretion onto a compact star in the Schwarzschild metric, which predicts periodic variations of luminosity when matter is supplied to the system at a constant accretion rate. These are related to the periodic appearance and disappearance of the Eddington Capture Sphere. In the model the frequency is found to vary inversely with the luminosity. If the input accretion rate varies (strictly) periodically, the luminosity variation is quasi-periodic, and the quality factor is inversely proportional to the relative amplitude of mass accretion fluctuations, with its largest value $Q\\approx 1/(10\\, |\\delta \\dot M/\\dot M|)$ attained in oscillations at about 1 to 2 kHz frequencies for a $2M_\\odot$ star. } ",
        "introduction": "\\cite{Abramowicz} (hereafter AEL) argued that the luminosity of a relativistic star that is accreting at a super-Eddington rate, should periodically change. They have shown that in the combined gravitational and radiation fields of a~spherical, compact star, radially moving test particles are captured by a sphere on which the gravitational and radiative forces balance. This is because in Einstein's general relativity the radiative force diminishes more strongly with the distance than the gravitational force, and radiation may be super-Eddington close to the star, but sub-Eddington further away, reaching the Eddington value at the \\emph{Eddington Capture Sphere} (ECS), whose radius is given by a simple expression in the Schwarzschild coordinates (\\cite{Phinney}), \\begin{equation} \\label{Eddington-radius} r_{\\rm ECS} =  \\frac{2 R_G}{1 - \\left(1 - \\dfrac{2 R_G}{R^2} \\right)^2 \\left(\\dfrac{L}{L_{\\rm Edd}} \\right)^2}\\ . \\end{equation} Here $R$ is the radius of the star, $R_G = GM/c^2$ its gravitational radius, and $L$, which is assumed to satisfy Eq.~(\\ref{cond}), is the stellar luminosity at its surface. \\cite{Bini,Oh}, and \\cite{Stahl} have shown that the ECS captures particles from a wide class of non-radial orbits as well\\footnote{As shown in detail by \\cite{Oh}, azimuthal radiation drag efficiently removes particle's angular momentum, typically making motions in the combined gravitational and radiation fields asymptotically radial.}, so the following discussion is not restricted to the case of radial accretion. However, it is important to keep in mind that the balance of forces necessary for the existence of the ECS requires a large, radial radiative flux. In non-spherical accretion this can only be realized in the optically thin regime.  In this paper we present a simple model in which the idea of periodic luminosity changes is realized. AEL assumed that the radiation power was provided by the kinetic energy of the accreted particles, released when they hit the surface of the star.  As stressed by \\cite{Stahl}, particles settle on the ECS rather gently, so the ECS itself is not particularly luminous. If all particles are captured at the Eddington sphere, they do not reach the surface of the star, and the stellar accretion luminosity goes to zero. When it does, and actually already when ${L}/{L_{\\rm Edd}} < (1-2R_G/R)^{-1/2}$, the ECS disappears, accretion is resumed by the star and eventually the luminosity may build up to its former value, so the ECS will reappear and accretion will stop again. Thus, the accretion process may be quasiperiodic, alternating between states of high luminosity and no luminosity. ",
        "conclusions": "\\label{discuss}  We have shown that supplying mass to the vicinity of a compact star, continuously and at a constant accretion rate, can lead to  periodic top-hat variations of luminosity, if only the mass accretion rate corresponds to a mildly super-Eddington luminosity at the stellar surface. This oscillatory behaviour of luminosity is related to the phenomenon of the Eddington Capture Sphere, which is a consequence of  the interplay of radiation drag and general relativity.  If such an oscillation occurs in the real world, for example in the Z sources, where rapid variations of the inferred inner radius of the accretion disk have been reported (\\cite{Lin}), the actual mechanism is likely to be more complex than the strictly periodic oscillation in the toy model considered here. E.g., the accretion rate is not likely to be constant (contrary to our assumption of simple on-off behaviour at the stellar surface).  As a minor extension of the model, consider an accretion rate that is varying on a timescale comparable to the period of the ECS oscillator. Even strictly periodic variation of $\\dot M_0$ will lead to a decoherence of the ECS oscillation, because the position $r_0$ of the ECS and (hence) the infall time from the ECS, and (hence) the oscillator frequency, are strongly varying functions of the luminosity, $L_0=\\eta M_0$. Indeed, \\begin{equation} \\label{decohere} \\frac{|\\delta f|}{f} = \\frac{|\\delta T|}{T} = \\frac{d\\ln T}{d\\ln L_\\infty}\\frac{|\\delta  L_\\infty|}{ L_\\infty} =\\frac{d\\ln T}{d\\ln L_\\infty}\\left|\\frac{\\delta \\dot M_0}{\\dot M_0}\\right|. \\end{equation} Thus, the $Q$ factor of the oscillation is \\begin{equation} \\label{Q} Q=\\frac{f}{|\\delta f|} = \\left|\\frac{\\dot M_0}{\\delta \\dot M_0}\\right| \\left(\\frac{d\\ln T}{d\\ln L_\\infty}\\right)^{-1}. \\end{equation} We have indicated some values of the logarithmic derivative ${d\\ln T}/{d\\ln L_\\infty}$ in Figs.~\\ref{fig:relation},~\\ref{fig:mass}. Comparing the figures, we see that delayed emission (reaction time $\\delta t_r>0$) at the stellar surface acts as a low-pass filter, cutting out the highest frequencies. At the same time it improves the quality factor of the oscillator.  Accreting neutrons stars in low mass X-ray binaries are expected to have a radius in the range of $4$ to $7 GM/c^2$ (\\cite{arnett,kw}), and a mass of about two solar masses. It is seen from Fig.~\\ref{fig:mass} that the minimum values of the logarithmic derivative attained for radii of 4, 5, 6, $7 GM/c^2$ are between 5 and 15. Hence, the maximum expected value of the oscillator quality factor would be $Q\\approx 1/(10 |\\delta\\dot M_0/\\dot M_0|)$, occurring for frequencies between about 1 and 2 kHz, as can be seen from Fig.~\\ref{fig:mass}."
    },
    "1208/1208.2422_arXiv.txt": {
        "abstract": "BH-NS and BH-BH systems are among the most promising gravitational wave sources detectable by advanced LIGO/VIRGO and the Einstein Telescope.  Although the rates of these systems may be above those of NS-NS mergers, BH-NS and BH-BH systems are difficult to detect, and thusfar none have been observed.  But the progenitors of BH-NS and BH-BH binary systems may have been observed, in the form of High-Mass X-ray Binaries (HMXBs).  In this paper, we continue work studying these potential progenitors of these important gravitational wave sources.  In the first two papers of the series, we have demonstrated that IC10 X-1 and NGC300 X-1 are direct progenitors of BH-BH systems and that Cyg X-1 may form, alas with a very low probability, a BH-NS system. Here, we analyze the Galactic binaries GX 301-2, Vela X-1, XTEJ1855-026, 4U1907+09, Cir X-1 and extra-galactic LMC X-1, LMC X-3, M33 X-1. In each case, we find that the future evolution will not allow the formation of a BH-NS nor a BH-BH system. Most of these binaries will soon merge in the common envelope phase, with a compact object sinking into a helium-rich core of a stellar companion.  This ``helium-merger'' may be a progenitor for long duration gamma-ray bursts (GRBs).  Based on the observed HMXB population, the rate of helium-mergers may make up a sizable fraction of long-duration GRBs. Due to this high number of potential GRB progenitors, a chance that a Galactic HMXB has caused one of the recent major mass extinction events is significant ($\\sim 10-20\\%$). ",
        "introduction": "Binary compact mergers are the most promising sources for ground-based gravitational wave detectors such as LIGO or VIRGO.  These observatories are currently being upgraded and are scheduled to begin operating in the 2015 time frame.  With the heightened sensitivity of these upgraded detectors, the questions surrounding gravitational wave detection are evolving from ``will we detect gravitational waves?'' to ``when will we detect gravitational waves?'' and ``what type of merging remnant will dominate the detections?''.  So far, astronomers have only observed double neutron stars (NS-NS), detecting the pulsar in radio frequencies of electromagnetic spectrum (e.g., Lorimer 2008). Kim, Kalogera \\& Lorimer (2010) assessed empirical NS-NS inpiral rates and it appears that the detection of these events in gravitational radiation is unavoidable once the advanced instruments reach their design sensitivity.  But the upgraded LIGO and VIRGO systems should also be sensitive to other types of double compact objects, namely black hole neutron star systems (BH-NS) and double black holes (BH-BH).  Unlike NS-NS binaries, these binary systems have yet to be observed.  In principle, these types of double compact objects are within reach of radio and microlensing surveys. Dong et al. (2007) reported the first potential detection of BH-BH system, but allowed for an alternative interpretation of the {\\tt OGLE-2005-SMC-001} lensing event.  Population synthesis results based on theoretical evolutionary considerations indicate that BH-NS and BH-BH systems are expected to form and populate the local universe (e.g., Lipunov, Postnov \\& Prokhorov 1997; Bethe \\& Brown 1998; De Donder \\& Vanbeveren 1998; Bloom, Sigurdsson \\& Pols 1999; Fryer, Woosley \\& Hartmann 1999; Nelemans, Yungelson \\& Portegies Zwart 2001, Belczynski, Kalogera \\& Bulik 2002, Voss \\& Tauris 2003). In particular, it was recently argued that BH-BH systems will dominate the gravitational radiation inspirals (Belczynski et al. 2010b).  At {\\em Warsaw Observatory} we have undertaken a program to provide empirical constraints on the existence of these undetected classes of compact binary systems.  Bulik, Belczynski \\& Prestwich (2011) studied the future evolution of two extra-galactic X-ray binaries hosting massive black holes with Wolf-Rayet companions: IC10 X-1 and NGC300 X-1. It was argued that these binaries will produce massive BH-BH systems. Based on a simple rate estimate, it was shown that the detection of such massive BH-BH inspirals by gravitational detectors is likely even in the existing data of the initial LIGO and VIRGO.  Belczynski, Bulik \\& Bailyn (2011) have studied the famous Galactic binary Cyg X-1 hosting a massive black hole and an O type companion. It was found that the most likely fate of this binary is disruption during the supernova explosion that forms a neutron star out of the O type companion. Cyg X-1 has only a small chance of forming a close BH-NS system. The rate of BH-NS systems inferred from Cyg X-1 is too low to make them likely sources for advanced gravitational radiation detectors.  However, it was noted that this finding represents only a lower limit on the detection of BH-NS systems as these objects may form along channels that do not involve Cyg X-1 like stage.  In this study we discuss the remaining Galactic and extra-galactic high mass X-ray binaries (HMXBs) that are putative progenitors for the BH-NS or BH-BH systems. The binaries under consideration host already one compact object, a neutron star or a black hole, and a massive companion that can in principle form the second compact object. We chose only those binaries that have rather well established system parameters. We do not have any strict definition for our sample as we have chosen (rather subjectively) binaries that either appear to have a chance to form BH-NS or BH-BH system, or were brought to our attention by someone from the compact object community. In some cases, we had to perform detailed calculations to establish the fate of a given system. However, in other cases it was enough to search literature for a better parameter estimate or just to reiterate the previously presented arguments to confirm already proposed conclusions.  As we will show, it appears that besides considered earlier binaries IC10 X-1, NGC300 X-1 and Cyg X-1, none other is likely to form BH-NS nor BH-BH system.  In \\S\\,2, we review the fates of each system individually.  In many cases, the binary system is expected to merge.  This merger of a compact remnant with a massive companion has been proposed as a progenitor of long-duration GRBs (Fryer \\& Woosley 1998).   We discuss the implications of these systems on the GRB population in \\S\\,3. ",
        "conclusions": "Including our previous studies we have shown that out of $11$ very massive HMXB with reasonably known binary parameters: $2$ (IC10 X-1, NGC300 X-1) will form close BH-BH systems, $1$ (Cyg X-1) will have a $30\\%$ chance to form a BH-NS system (with $1\\%$ probability of this being a close system), $1$ (LMC X-3) will form a wide BH-WD system, and the remaining $7$ binaries will merge in common envelope (GX 301-2, Vela X-1, XTE J1855-026, 4U 1907+09, Cir X-1, LMC X-1, M33 X-1). The summary of our results is given in Table~1.  In the helium merger model of long-duration gamma-ray bursts (GRBs) the compact object, either a NS or a BH, sinks into the helium rich core of its companion. A compact object accretes at a high rate $\\gtrsim 0.01 \\mps$ (Fryer \\& Woosley 1998, Zhang \\& Fryer 2001) and, in the NS case, a BH forms. The helium core as it accretes onto a BH forms a transient torus (Barkov \\& Komissarov 2011) leading to a standard scenario for a GRB engine (Popham, Woosley \\& Fryer 1999).  Helium merger configuration is most easily provided in the common envelope merger of HMXBs in which an optical companion is already an evolved star beyond main sequence (He core) or at the end of main sequence (high concentration of He in the core). In the merger, part of a companion star is ejected (but not fully) away from the system, and the compact object sinks into helium-rich core.  The so-called ``Christmas burst'' (GRB 121225) may be potentially the first identifiable example of a helium merger GRB, with indication of an inner compact source and a surrounding shell (Thone et al. 2011).  Fryer, Belczynski \\& Thone (2013) have employed theoretical population synthesis predictions for common envelope mergers and estimated the characteristic observational features of helium merger GRBs. Their model calculates a GRB luminosity based on a helium core mass (required to be typically larger than $4 \\msun$) and establishes the position of the shell surrounding the central GRB engine based on the radius of the companion star (CE donor). Here we use our empirically based estimates and employ Fryer et al.  (2013) model to estimate typical luminosities of helium merger GRBs.  In the case of GX 301-2, we expect a NS to sink into a $60 \\rsun$ Hertzsprung gap star with $15 \\msun$ helium core and GRB luminosity of $L_{\\rm grb} = 10^{49}, 10^{51}$ erg s$^{-1}$ for neutrino annihilation and Blandford-Znajek emission model respectively (for each system, we will provide two values for $L_{\\rm grb}$ corresponding to these two emission mechanisms). For Vela X-1 a NS will enter a $32 \\rsun$ Hertzsprung gap star with $7 \\msun$ helium core and  $L_{\\rm grb} = 10^{46}, 10^{50}$ erg s$^{-1}$. In binary XTE J1855-026 a NS will merge into $26 \\rsun$ Hertzsprung gap star with $7 \\msun$ helium core and $L_{\\rm grb} = 10^{46}, 10^{50}$ erg s$^{-1}$. In case of 4U 1907+09 we predict a merger of a NS and a $26 \\rsun$ late main sequence or early Hertzsprung gap star with $8 \\msun$ helium-rich core and $L_{\\rm grb} = 10^{47}, 5 \\times 10^{50}$ erg s$^{-1}$. For Cir X-1 a NS will sink into a $12-19 \\rsun$ Hertzsprung gap star with $0.6-2.1 \\msun$ helium core.  For this case, a GRB is not expected but such mergers may account for some subset of ultraluminous supernovae (Fryer et al. 2013). For LMC X-1 a BH will merge into a $\\sim 11 \\msun$ helium-rich core of a $17.3 \\rsun$ star at the end of its main sequence evolution with a potential GRB at $L_{\\rm grb} = 10^{48}, 10^{51}$ erg s$^{-1}$. A BH in M33 X-1 will sink into a $21.3 \\rsun$ main sequence companion that has evolved through about $60-70\\%$ of its central hydrogen burning. Since the optical companion has mass of $\\sim 70\\msun$ its helium-rich core is also very massive $\\sim 20 \\msun$ and potential GRB luminosity is very high $L_{\\rm grb} = 10^{51}, 5 \\times 10^{51}$ erg s$^{-1}$.  There are at least 4 HMXB binaries in our Galaxy with very massive donors ($M_{\\rm opt}>20 \\msun$; required for a GRB) that will merge in the soon-to-come common envelope phase (see Tab.~1). As discussed above the merger will be followed shortly by a gamma ray burst (Fryer \\& Woosley 1998; Thone et al. 2011). The HMXB lifetime may be estimated from the lifetime of an optical component. For the considered binaries (GX 301-2, Vela X-1, XTE J1855-026, 4U 1907+09) the HMXB lifetime can be roughly estimated as the difference between the lifetime of the companion star and the minimum time required to form a NS i.e. $5$ Myr (Belczynski \\& Taam 2008).  The companion lifetime depends on its mass and it is $5.3$ Myr for a $43\\msun$ star (GX 301-2), $8$ Myr for $24\\msun$ star (Vela X-1), $7.8$ Myr for $25\\msun$ star (XTE J1855-026), and $7.3$ Myr for $27 \\msun$ star (4U 1907+09). This leads to the X-ray lifetime estimates of $\\tau=$ $0.3$, $3$, $2.8$ and $2$ Myr, respectively. The expected formation rate of such X-ray binaries is $R = \\sum \\tau_i^{-1}=4.5$ Myr$^{-1}$. Since the HMXB lifetimes are very short, this is also the estimate of the helium merger Galactic GRB rate. The rate estimate is formally dominated by the contribution of GX 301-2 and in this case this system deserves a few more comments. The companion in GX 301-2 might have gained up to $20 \\msun$ from the NS progenitor, and the rejuvenation might not have been complete. Thus it is quite likely that the expected lifetime of GX 301-2 as an HMXB is longer, up to $3$ Myr. If we adopt such long lifetime the formation rate of helium merger GRBs drops to $1.5$ Myr$^{-1}$.  These rates have been calculated assuming that the entire Galaxy has been searched for HMXB and our observational sample is complete. In reality HMXBs have been found only in about $25\\%$ of the volume of the disk (e.g., Ozel et al. 2010). Thus the estimated rates have to be corrected for that and are in the range $R_{\\rm Galactic}=6-18$ Myr$^{-1}$.  The star formation in the Galactic disk is approximately constant and at the level of $3.5 \\mpy$ (e.g., O'Shaughnessy et al. 2006). The rate of helium merger GRBs may be translated into ${\\cal R}_{grb} \\approx 1.6 - 5.1 \\times 10^{-6} \\msun^{-1}$ and is expressed per unit of star forming mass. This value can be used to estimate the GRB rate in other galaxies or in cosmological population studies. Since the delay of helium CE merger events is rather short ($\\lesssim 10$ Myr), these particular GRBs are expected in any galaxy with ongoing star formation.  It must be noted that our estimate provides only a lower limit on the helium merger GRB rate. There are a number of HMXBs in our Galaxy without fully established orbital/component parameters (Liu, van Paradijs \\& van den Heuvel 2006), some of which may be potential GRB progenitors. There may be ways to produce a helium merger GRB without a HMXB phase, for example NS ejection into a companion via supernova natal kick (e.g., Fryer et al. 2013).  Nevertheless, our estimate provides a rather significant number of GRBs in the history of the Galactic disk $\\sim 60,000-180,000$ (for $10$ Gyr of constant star formation) and we check whether such a GRB may have affected life on Earth. The extent of the disk, where most of star formation is taking place, may be approximated by radius of $18$ kpc and vertical extent of $0.60$ kpc (Paczynski 1990) and that gives disk volume of $V_{\\rm disk} \\approx 610$ kpc$^3$. Long GRBs with their maximum energy output of $5 \\times 10^{51}$ erg/s lasting $\\sim 10$s may have a big impact or are claimed to be lethal for Earth biosphere if exploding within $2$ kpc (Nakar 2010). The volume of impact sphere within the disk is then about $V_{\\rm impact} \\approx 7.5$ kpc$^3$. And therefore the Galactic rate needs to be reduced by factor of $\\sim 80$ ($V_{\\rm disk}/V_{\\rm impact}$) to provide the rate estimate of GRBs in the impact zone: $0.073-0.22$ Myr$^{-1}$. The long GRB power is collimated in a relativistic jet with beaming of $\\sim 5^{\\circ}$ (e.g., Frail et al. 2001; Soderberg et al. 2006). Since only a direct hit on the Earth biosphere may be lethal or cause a major mass extinction, this leads to further reduction by factor of $\\sim 500$. Finally, we obtain the rate of direct GRB hits coming from the impact zone at the level $0.15-0.45$ Gyr$^{-1}$. Over a $4.5$ Gyr of Earth history the expected number of direct nearby hits is in the range  $0.65 - 2.0$. There are at least five major mass extinction events noted in the last $0.5$ Gyr in the marine data (Alroy 2008). The probability of direct impact from  a helium merger GRBs in this period is in the range $\\sim 7.2 - 22\\%$). Thus it is quite possible that one of these events has been caused by a nearby helium GRB."
    },
    "1208/1208.1617_arXiv.txt": {
        "abstract": "Foreground removal is a major challenge for detecting the redshifted $21$-cm neutral hydrogen (HI) signal from the Epoch of Reionization (EoR). We have used $ 150 \\, {\\rm MHz}$ GMRT observations to characterize the statistical properties of the foregrounds in four different fields of view.  The measured multi-frequency angular power spectrum $C_{\\ell}(\\Delta \\nu)$ is found to have values in the range $10^4 \\ {\\rm mK^2}$ to $2 \\times 10^4 \\ {\\rm mK}^2$ across $700 \\le \\ell \\le 2 \\times 10^4$ and $\\Delta \\nu \\le 2.5 \\ {\\rm MHz}$, which is consistent with model predictions where point sources are the most dominant foreground component.  The measured $C_{\\ell}(\\Delta \\nu)$ does not show a smooth $\\Delta \\nu$ dependence, which poses a severe difficulty for foreground removal using polynomial fitting.  The observational data was used to assess point source subtraction. Considering the brightest source $(\\sim 1 \\ {\\rm Jy})$ in each field, we find that the residual artifacts are less than $1.5 \\%$ in the most sensitive field (FIELD I). Considering all the sources in the fields, we find that the bulk of the image is free of artifacts, the artifacts being localized to the vicinity of the brightest sources. We have used FIELD I, which has a rms noise of $1.3 \\ {\\rm mJy/Beam}$, to study the properties of the radio source population to a limiting flux of $9 \\ {\\rm mJy}$. The differential source count is well fitted with a single power law of slope $-1.6$.  We find there is no evidence for flattening of the source counts towards lower flux densities which suggests that source population is dominated by the classical radio-loud Active Galactic Nucleus (AGN).  The diffuse Galactic emission is revealed after the point sources are subtracted out from FIELD I . We find $C_{\\ell} \\propto \\ell^{-2.34}$ for $253 \\le \\ell \\le 800$ which is characteristic of the Galactic synchrotron radiation measured at higher frequencies and larger angular scales.  We estimate the fluctuations in the Galactic synchrotron emission to be $ \\sqrt{\\ell(\\ell+1)C_{\\ell}/2\\pi} \\simeq 10 \\, {\\rm K}$ at $\\ell=800$ ($\\theta > 10'$). The measured $C_{\\ell}$ is dominated by the residual point sources and artifacts at smaller angular scales where $C_{\\ell} \\sim 10^3 \\ {\\rm mK}^2$ for $\\ell > 800$. ",
        "introduction": "Observations of the high-redshift Universe with the 21-cm hyperfine line emitted by neutral hydrogen gas (HI) during the Epoch of Reionization (EoR) contains a wealth of cosmological and astrophysical informations (see \\citealt{furla}; \\citealt{morales10} for recent reviews). Analysis of quasar absorption spectra \\citep{becker,fan} and the CMBR \\citep{komatsu} together indicate that reionization probably started around a redshift of $15$ and lasted until a redshift of $6$. The Giant Metrewave Radio Telescope (GMRT \\footnote{http://www.gmrt.ncra.tifr.res.in}; \\citealt{swarup}) is currently functioning at several bands in the frequency range of $150 \\, {\\rm MHz}$ to $1420 \\, {\\rm MHz}$ and can potentially detect the 21 cm signal at high redshifts.  Several upcoming low-frequency telescopes such as LOFAR\\footnote{http://www.lofar.org/}, MWA\\footnote{http://www.mwatelescope.org/} and PAPER\\footnote{http://astro.berkeley.edu/\\textasciitilde dbacker/eor/} are also targeting a detection of the high redshift 21-cm signal. LOFAR is already operational and will soon start to collect data for the EoR project.  In this paper we have carried out GMRT observations to characterize the statistical properties of the background radiation at $150 \\, {\\rm MHz}$.  These observations cover a frequency band of $\\sim 5 \\ {\\rm MHz}$ with a $\\sim 100 \\ {\\rm kHz}$ resolution, and cover the angular scales $\\sim 30'$ to $\\sim 30''$.  The 21 cm radiation is expected to have an rms brightness temperature of a few mK on angular scales of a few arc minute \\citep{zald}. This signal, is however, buried in the emission from other astrophysical sources which are collectively referred to as foregrounds.  Foreground models \\citep{ali} predict that they are dominated by the extragalactic radio sources and the diffuse synchrotron radiation from our own Galaxy which makes a smaller contribution. Analytic estimates of the HI 21-cm signal indicate that the predicted signal decorrelates rapidly with increasing $\\Delta\\nu$ and the signal falls off $90\\%$ or more at $\\Delta\\nu \\ge 0.5 \\ {\\rm MHz}$ \\citep{BS1,BA5}. The foreground from different astrophysical sources are expected to be correlated over a large frequency frequency separation. This property of the signal holds the promise of allowing us to separate the signal from the foregrounds \\citep{ghosh2}. In this paper we have characterized the foreground properties across a frequency separation $\\Delta \\nu$ of $2.5 \\ {\\rm MHz}$, the signal may be safely assumed to have decorrelated well within this frequency range.  There currently exist several surveys which cover large regions of the sky at around $150 \\,{\\rm MHz}$ like the 3CR survey \\citep{edge,Bennett}, the 6C survey \\citep*{hales,waldram} and the 7C survey \\citep{hales07}.  These surveys have a relatively poor angular resolution and a limiting flux density of $\\sim 100\\, {\\rm mJy}$. \\citet{dmat1} have used the 6C survey, the 3CR survey and the 3 CRR catalogue \\citep{laing} to estimate the extragalactic point source contribution at $150 \\, {\\rm MHz}$.  The GMRT, which is currently the most sensitive telescope at $150 \\ {\\rm MHz}$, offers the possibility to perform deeper surveys with a higher angular resolution.  A few single pointing $150 \\, {\\rm MHz}$ surveys \\citep{sirothia,ishwara10,ishwara,intema} have been performed with an rms noise of $\\sim 1 \\ {\\rm mJy/Beam}$ at a resolution of $20''-30''$. The TGSS \\footnote{http://tgss.ncra.tifr.res.in/} is currently underway to survey the sky north of the declination $-30^{\\circ}$ with an rms noise of $7-9 \\ {\\rm mJy/Beam}$ at an angular resolution of $20''$.  Diffuse synchrotron radiation produced by cosmic ray electrons propagating in the magnetic field of our Galaxy \\citep{GS69} is a major foreground component. This component, which is associated with our Galactic disk, is particularly strong in the plane of our Galaxy. Our target fields were selected well off the Galactic plane with $b>10^{\\circ}$ where we expect a relatively lower contribution from the Galactic synchrotron radiation.  Radio surveys at 408 MHz \\citep{haslam}, 1.42 GHz \\citep{R82,RR88} and 2.326 GHz \\citep{JBN98} have measured the diffuse Galactic synchrotron radiation at angular scales larger than $\\sim 1^{\\circ}$ where it is the most dominant contribution, exceeding the point sources. \\citet{platania1} have analyzed these observations to show that the Galactic synchrotron emission has a steep spectral index of ${\\alpha} \\approx 2.8$.  The angular power spectrum $C_{\\ell}$ has been measured at $2.3 \\,{\\rm GHz}$ \\citep{giardino01} and $2.4 \\,{\\rm GHz}$ \\citep{giardino02} where they find $C_{\\ell} \\sim \\ell^{-2.4}$ at angular multipoles $\\ell < 250$. WMAP data shows $C_{\\ell} \\sim \\ell^{-2}$ at $\\ell \\le 200$ \\citep{bennett03} which is slightly flatter compared to the findings at lower frequencies.  It is relevant to note that all of these results are restricted to angular scales greater than $43'$ $(\\ell < 250)$, little is known (except \\citealt{bernardi09} discussed later) about the structure of the Galactic synchrotron radiation at the sub-degree angular scales probed in this paper.  \\citet{ali} have carried out GMRT observations to characterize the foregrounds on sub-degree angular scales at $150 \\ {\\rm MHz}$. They have used the measured visibilities to directly determine the multi-frequency angular power spectrum $C_{\\ell}(\\Delta \\nu)$ (MAPS; \\citet{kanan}) which characterizes the statistical properties of the fluctuations in the background radiation jointly as a function of the angular multipole $\\ell$ and frequency separation $\\Delta \\nu$. They find that the measured $C_{\\ell}(\\Delta \\nu)$ has a value of around $10^4 \\ {\\rm mK}^2$ which is seven orders of magnitude larger than the expected HI signal. The measured $C_{\\ell}(\\Delta \\nu)$ was also found to exhibit relatively large ($\\sim 10 - 20 \\%$) fluctuations in $\\Delta \\nu$ which pose a severe problem for foreground removal.  Further, it was attempted to remove the foregrounds by subtracting out all the identifiable point sources above $8 \\, {\\rm mJy}$. The resulting $C_{\\ell}$, however, only dropped to $\\sim 3$ times the original value due to residual artifacts.  \\citet{bernardi09} have analyzed $150 \\, {\\rm MHz}$ WSRT observations near the Galactic plane ($(l=137^{\\circ}, b=+8^{\\circ}$) to characterize the statistical properties of the diffuse Galactic emission (after removing the point sources) at angular multipoles $\\ell < 3000$ ($\\theta > 3.6'$).  They find that the measured total intensity angular power spectrum shows a power law behaviour $C_{\\ell} \\propto \\ell^{-2.2}$ for $\\ell \\le 900$.  They also measured the polarization angular power spectrum for which they find $C_{\\ell} \\propto \\ell^{-1.65}$ for $\\ell \\le 2700$. Their observations indicate fluctuations of the order of $\\sim 5.7 \\ {\\rm K}$ and $3.3 \\ {\\rm K}$ on $5'$ angular scales for the total intensity and the polarized diffuse Galactic emission respectively.  \\citet{pen09} have carried out $150 \\, {\\rm MHz}$ GMRT observations at a high Galactic latitude to place an upper limit of $\\sqrt{\\ell^2C_{\\ell}/2\\pi} < 3 {\\rm K}$ on the polarized foregrounds at $\\ell < 1000$.  There has been a considerable amount of work towards simulating the foregrounds \\citep{jelic,sun,bowman09,gleser,harker,Liu1,liu09b,nada} in order to develop algorithms to subtract the foregrounds and detect the redshifted 21-cm signal. These simulations require guidance from observational data, and it is crucial to accurately characterize the foregrounds in order to have realistic simulations for future observations.  One of the challenges for the detection of this cosmological 21-cm signal is the accuracy of the foreground source removal. Our current paper aims at characterizing the foreground at arcminute angular scales which will be essential for extracting the 21-cm signal from the data. In this paper we study the radio source population at $150 \\, {\\rm MHz}$, which is still less explored, to determine the differential source count which plays a very important role in predicting the angular power spectrum of point sources \\citep{ali,ghosh1}. We find that the point sources are the most dominating foreground component at the angular scales of our analysis. The accuracy of point source removal is one of the principal challenges for the detection of the 21-cm signal. We investigate in detail the extragalactic point source contamination in our observed fields and study how accurately the bright point sources can be removed from our images. We also estimate the level of residual contamination in our images. It is expected that the diffuse Galactic emission will be revealed if the point sources are individually modelled and removed with high level of precession. To our knowledge, the present work is only the second time that the fluctuations in the Galactic diffuse emission have been detected at around $10$ arcminute angular scales at a frequency of $150 \\, {\\rm MHz}$ .  We also complement this with a measurement of the angular power spectrum of the diffuse Galactic foreground emission in the $\\ell$ range $200$ to $800$.  In this paper we have analyzed GMRT observations to characterize the foregrounds in four different fields at $150 \\, {\\rm MHz}$ which corresponds to the HI signal from $z = 8.3$. We note that unless otherwise stated we use the cosmological parameters $(\\Omega_{m0},\\Omega_{\\Lambda0}, \\Omega_b h^2, h, n_s,\\sigma_8)=(0.3,0.7, 0.024,0.7,1.0,1.0)$ in our analysis. We now present a brief outline of the paper. Section 2 describes the observations and data analysis.  We have used the multi-frequency angular power spectrum $C_{\\ell}(\\Delta \\nu)$ to quantify the statistical properties of the background radiation in the observed fields. Section 3 presents the technique that we have used to estimate $C_{\\ell}(\\Delta \\nu)$ directly from the measured visibilities.  We show the the measured $C_{\\ell}(\\Delta \\nu)$, and discuss its properties in Section 4. In this section we have also compared the measured $C_{\\ell}(\\Delta \\nu)$ with foreground model predictions.  In Section 5 we consider the extragalactic point sources which forms the most dominant foreground component in our observations. In order to assess how well we can subtract out the point sources, we consider point source subtraction for the brightest source in each field, as well as subtracting out all the identifiable sources.  We use the most sensitive field (FIELD I) in our observation to study the nature of the radio source population at $150 \\, {\\rm MHz}$, and we also provide a source catalogue for this field in the online version of this paper (Appendix \\ref{catal}).  In Section 6 we have used the residual data of FIELD I, after point source removal, to study the diffuse Galactic synchrotron radiation. We have summarized the results of the entire paper in Section 7. ",
        "conclusions": "We have analyzed $150 \\, {\\rm MHz}$ GMRT observations in four different fields of view (Table~\\ref{tab:obs_sum}, Figure \\ref{fig:fields}). We have used the multi-frequency angular power spectrum (MAPS) $C_{\\ell}(\\Delta \\nu)$ to jointly characterize the statistical properties of the fluctuations as a function of the angular multipole $\\ell$ and the frequency separation $\\Delta\\nu$ (Figure \\ref{fig:clcont}) across the range $700 \\le \\ell \\le 2 \\times 10^4$ and $0 \\le \\Delta \\nu \\le 2.5 \\, {\\rm MHz}$.  We find that the measured $C_{\\ell}$ has values around $10^4 \\, {\\rm mK^2}$ at $\\ell \\sim 1000$, and it drops by around $50 \\%$ at $\\ell\\sim 10^4$. The measured $C_{\\ell}$ is foreground dominated and is more than $7$ orders of magnitude larger than the expected HI signal.  Analytic estimates of the HI signal \\citep{BA5} show that we expect $C^{HI}_{\\ell}(\\Delta \\nu)$ to decorrelate rapidly with increasing $\\Delta \\nu$, and fall by $90\\%$ or more at $\\Delta\\nu \\ge 0.5 \\ {\\rm MHz}$ . In contrast, the foregrounds from different astrophysical sources are expected to be correlated over large frequency separations ($\\Delta \\nu > 2 \\, {\\rm MHz}$). This holds the promise of allowing us to separate the signal from the foregrounds \\citep{ghosh2}. We find that the measured $C_{\\ell}(\\Delta \\nu)$ (for reference see Figures \\ref{fig:clnuf1} and \\ref{fig:clnuf2}) shows a smooth $\\Delta \\nu$ dependence with a $20$ to $40 \\, \\%$ variation across the measured $2.5 \\, {\\rm MHz}$ $\\Delta \\nu$ range. However, contrary to our expectations, in addition to a smooth behaviour we also notice abrupt variations and oscillations $(\\le 10 \\%)$ in the $\\Delta \\nu$ dependence of the measured $C_{\\ell}(\\Delta \\nu)$. These abrupt variations and oscillations, whose origin is at present not understood, poses a severe impediment for foreground removal. Polynomial fitting along the frequency axis has been extensively considered for foreground removal.  All attempts in this direction may fail due to the presence of abrupt variations and small oscillations which cannot be filtered out with low order polynomials.   The measured $C_{\\ell}$'s, we find, are consistent with foreground models which predict that extragalactic point sources make the most dominant contribution at the sub-degree scales that we have probed here (Figure \\ref{fig:clfg}).  The brightest source ($\\sim 1 \\, {\\rm Jy}$) in each field alone contributes around $10 \\%$ of the total measured $C_{\\ell}$.  It is very important to correctly identify the point sources and subtract these out at a high level of precision. We have carried out point source subtraction in all the four fields that we have analyzed here (Figures \\ref{fig:bs1} and \\ref{fig:subfields}). Considering only the brightest source, we find that source subtraction is most effective in FIELD I, where the image also has the lowest rms noise of $1.3 \\, {\\rm mJy/Beam}$ (Table~\\ref{tab:img_sum}). The residual artifacts after subtracting out the brightest source, we find, are within $1.5 \\%$ of the of the original source flux ($905 \\, {\\rm mJy}$).  The subsequent analysis was restricted to FIELD I which has the lowest rms noise. The Clean Components corresponding to all the sources above $10 \\, {\\rm mJy}$ were subtracted out from the visibility data using the AIPS task UVSUB. We find that the resulting angular power spectrum $C_{\\ell}$ falls to $\\sim 1,000 \\, {\\rm mK}^2$ in the $\\ell$ range $800$ to $8,000$ (Figure \\ref{diffgalps}), which is roughly one-tenth of the $C_{\\ell}$ before source subtraction.  At these multipoles, we interpret the measured $C_{\\ell}$, after source subtraction, as arising from a combination of the residual artifacts from the bright sources and the faint sources ($< 10 \\, {\\rm mJy}$) that have not been removed.  The behaviour at lower multipoles ($\\ell \\le 800$), we find, is well fitted by a power law $C_{\\ell}= (513\\pm 41 ) \\times (1000/\\ell)^{2.34 \\pm 0.28} \\,{\\rm mK^2}$ which we interpret as the contribution from the diffuse Galactic synchrotron radiation.  The measured slope is consistent with earlier WSRT $150 \\, {\\rm MHz}$ observations \\citep{bernardi09}, and also with $2.3$ and $2.4 \\ {\\rm GHz}$ results at smaller $\\ell$ \\citep{giardino01,giardino02}, whereas WMAP finds a flatter slope $C_{\\ell} \\sim \\ell^{-2}$ \\citep{bennett03} at smaller $\\ell$ and much higher frequencies ($23-94 \\ {\\rm GHz}$)."
    },
    "1208/1208.3338_arXiv.txt": {
        "abstract": "{% Magnetic fields are regarded as a crucial element for our understanding of stellar physics. They can be studied with a variety of methods which provide complementary -- and sometimes contradictory -- information about the structure, strength and dynamics of the magnetic field and its role in the evolution of stars. Stellar magnetic fields can be investigated either with direct methods based on the Zeeman effect or through the observation of activity phenomena resulting from the interaction of the field with the stellar atmosphere. In this Cool Stars \\textsc{XVII} Splinter Session we discussed the results obtained by the many ongoing studies of stellar activity and direct studies of surface magnetic fields, as well as the state-of-the-art techniques on which they are based. We show the strengths and limitations of the various approaches currently used and to point out their evolution as well as the interest of coupling various magnetism and activity proxies. } ",
        "introduction": "Since the first detection on the Sun through the Zeeman effect in 1908 by G.~E. Hale, magnetic fields have become a key element of stellar physics. They are now thought to play an important role throughout the formation and evolution of stars and their planetary systems. In addition, they are now detected across a large part of the H-R diagram (e.g. Donati \\& Landstreet 2009; Reiners 2012), either directly or through the activity phenomena they power which can be observed across a large part of the electromagnetic spectrum. The characterization of these magnetic fields has revealed that stars show a very wide variety of magnetic properties in terms of strength, geometry or variability. Drawing a clear picture of stellar magnetism and dynamo processes is however not easy since different types of measurements do not always agree -- at least apparently -- and have led to various interpretations (e.g. Reiners et al. 2009; Morin et al. 2011 in the case of M dwarfs).  In this session, we discussed a number of approaches used to detect, measure and model stellar magnetic fields as well as the related activity phenomena. A particular emphasis was put on the comparison of different approaches or proxies, the way they can be combined, and on the ongoing developments in instrumentation, analysis and modelling. ",
        "conclusions": "\\label{sec:summary} The interpretation and comparison of magnetic field measurements and modelling stemming from different approaches is a debated topic. There is, however, a general agreement that a better knowledge of the spatial structure of stellar magnetic fields is essential in order to reconcile these different views, and that what we know about the Sun likely cannot be safely extrapolated to stars with very different properties (in particular very active stars). For instance, the predominance of bright (plage-like) \\textit{vs} dark (spot-like) regions among magnetic surface features, the very existence of regions corresponding to ``quiet'' photosphere on very active stars, are of great relevance to magnetic field measurements and start to be addressed through numerical simulations (sec.~\\ref{sec:bbeeck}). This importance of the spatial structure of the field is even relevant to the vocabulary used. It can be argued that ``magnetic field'' should be reserved for an actual magnetic field vector at some point of the stellar surface, whereas observationally-derived quantities should be termed ``magnetic flux''. Both quantities have the same unit of Gauss (CGS) or Tesla (SI), although some confusion can arise with a quantity also called ``magnetic flux'' and measured in Maxwell (i.e. G.cm$^2$, CGS) or Weber (i.e. T.m$^2$, SI), see Reiners (2011) and Reiners \\& Mohanty (2012).  The talks given in the session suggest that in the next few years progress will arise from the use of new or upgraded observational facilities, coupled with improved modelling methods and from multi-wavelength or multi-technique studies. Ultimately, such advances will aim at a better characterization of stellar magnetic fields and planetary environments, as well as a better understanding of magnetic field generation in cool stars."
    },
    "1208/1208.3754_arXiv.txt": {
        "abstract": " ",
        "introduction": "This document has a very specific audience. It is addressed to people who have, or who have been landed with, the responsibility for developing a \\gls{DMP} policy for a \\q{big science} collaboration, or some similar multi-institutional or multi-national project with a need for a bespoke plan.  Although it is nominally addressed to this (rather small) readership, we have written it with the intention that it will additionally be of use to: \\begin{itemize} \\item those evaluating or assessing such plans, for example within funders; and \\item people developing similar bespoke plans for scientific and other entities at this or other scales, who are looking for practical guidance on where to start, but for whom existing \\gls{DMP} guidance is too low-level or mechanical. \\end{itemize} For the purposes of this report, we presume that the reader is broadly persuaded (by external fiat if nothing else) of the need to preserve research data appropriately, and that they have both sophisticated technical support and the budget to support bespoke developments where necessary, obtained from a broadly supportive funder.  We take the position that: \\begin{itemize} \\item the demand for principled data management and data sharing is a reasonable one, and note that publicly funded projects typically have no fundamental objections to it; \\item that a reasonable framework for at least approaching the problem already exists in \\gls{OAIS} (\\prettyref{s:whyoais}); \\item that the OAIS recommendation is (just) concrete enough that it is not merely waffle; and \\item that there is a bounded set of resources which, if mastered by the reader, will allow them to produce a project DMP plan which is practically acceptable to the project, and discharges the principled demands of the funder and of society. \\end{itemize} Within this report we have sought to represent a consensus of views across the \\q{large-science} community within the UK, both through the roles of the authors of this document and also through a wider consultation we have undertaken with funders and research leaders. For more specific acknowlegements, see the section on~p.\\pageref{s:acknowledgements}.   The document is structured into three parts. \\begin{itemize} \\item \\prettyref{s:policy}, policy background: this part discusses the various high-level policy drivers for DMP planning. We take it as read that an organisation is aware of the need to manage its data professionally, in order that this data is readily accessible to the researchers within it. However, there are a number of higher-level interests which must be respected, concerning longer-term disciplinary goals, and the goals of society at large. \\item \\prettyref{s:technical}, technical background: this part is mostly about the technical frameworks relevant to the good management of data, and in particular the \\gls{OAIS} model.  We believe this is the key set of technologies which someone producing a project DMP policy should be aware of. \\item \\prettyref{s:planning}, DMP planning: everything more specific, which includes some discussion of the (poorly-modelled) costs of such preservation, and of existing work on validating (and its conjugate, auditing) DMP plans.  Though this section is more detailed than the earlier ones, it is not concerned with the nitty-gritty of RAID, network or NAS management, which are the province of the DMP plan's implementers. \\end{itemize}  \\q{Data management} does not contain many profound imponderables; navels need not be gazed at. Though it is going too far to say that we are peddling organised common sense, the majority of the relevant background material is readily accessible, as long as it can be found, and be known to be relevant. Our practical goal in this document is to assemble and contextualise this background material, arrange it in a way which is useful to the consituency we are aiming at, indicate where best practice may be found or where it is still unknown, and thereby enable the reader to lead the development of a DMP plan for their organisation, secure in the knowledge that they have a reasonable claim to be on top of the relevant literature.  \\subsection{Focuses, coverage, and some definitions}  The document is practical in tone, necessarily without being prescriptive; however, for our intended audience, the \\q{practical} includes some aspects of the larger policy background which must be respected, so we include coverage of these aspects, as well.  The report has been produced with a UK focus, but the only places where this is, we believe, apparent are in the UK emphasis of the policy discussion in \\prettyref{s:rcuk-principles}, and on the prominence of the \\gls{STFC} in our definition of big science below. Although STFC is (for this reason) particularly prominent, there is \\q{big science} data also to be found in research supported by the \\gls{EPSRC}, the \\gls{BBSRC} and the \\gls{NERC}.  There is more context available in the document \\q{Managing Research Data in Big Science}~\\cite{gray12}.  This is the final report of a project funded by JISC in 2010\\range 11, which was concerned with the background for big-science data management in general, and this present report in some places draws text directly from the earlier one. This might be useful for fuller discussion or further references, and we will make occasional reference to it in order to keep this present document short.  Throughout, the report is informed where appropriate by the OAIS reference model.  The model is introduced as technical background in \\prettyref{s:oais-description}, and more details are discussed in that section and as details of practice in \\prettyref{s:practice-validation}, but the ideas are pervasive enough that we feel it is useful to give a brief informal description of the model and its advantages at the beginning of the document, in \\prettyref{s:whyoais}.  For clarity, it seems useful to make briefly explicit what we mean by \\gls{DMP} and the term \\q{big science}, and we do this in the subsections below.  \\subsection{The what, why and how of OAIS} \\label{s:whyoais}  As suggested above, this document's advice orbits around the \\gls{OAIS} standard, adopting its (useful) concepts and vocabulary, and making reference to the other work on validation and costing that builds on it.  It is therefore useful to briefly discuss the \\q{what?}, \\q{why?} and \\q{how?} of OAIS, in that order.  \\emph{What is the OAIS model?} The OAIS reference model~\\cite{std:oais} is a conceptual model of the functions and responsibilities of an archive of (typically) digital objects, where the archive is viewed as an organisation or other entity, in principle distinct from the data producer, which exists to preserve those objects into the \\gls{Long Term}.  The OAIS standard does not describe how to achieve this, but it \\emph{does} clearly articulate the various steps of the process (for example that data goes through phases of Submission to an archive, Preservation there, and Dissemination to users), the various roles involved (for example data \\glspl{Producer} versus \\glspl{Consumer}), and what, at a high level, has to be done to let all this happen (for example the creation and management of documentation about \\gls{Representation Information}).  There is a fuller description of OAIS in \\prettyref{s:oais-description}.  \\emph{Why should you care?}  Integral to its development, the OAIS standard defines a fairly extensive vocabulary for digital preservation (each of the capitalised terms in the preceding paragraph has a precisely defined meaning, and when such terms appear below they are included in the glossary at the end), and although none of these definitions is particularly startling, and although the standard text can seem a little verbose, verging on windy, these terms have become the standard ones, and most work in this area is framed, directly or indirectly, by the OAIS concept set.  Thus, although the OAIS model is not the \\emph{only} model for a digital archive (see \\prettyref{s:dcc-lifecycle} for another), it is both plausible and conventional, and so makes a good starting point, and a useful shared understanding, for any discussion of digital preservation.  In addition, it is worth pointing out that the model was developed by the \\gls{CCSDS}, and so has a heritage which makes it a natural fit for non-space science data.  \\emph{How do I implement an OAIS model?} There is no general recipe, and by assumption the readers of this document are interested in systems which are large or unusual enough that no recipe is likely to be applicable.  Instead, we aim to provide pointers to resources which guide you in the right direction, and possibly reassure you that there are no major areas of concern you have missed. To start with, there is the brief introduction below in \\prettyref{s:oais-description}, plus tutorial reports such as~\\cite{lavoie04}, and book-length resources such as~\\cite{giaretta11}.  \\emph{OK, how do I know \\emph{when I have implemented} an OAIS model?} The OAIS model can be criticised for being so high-level that \\qq{almost any system capable of storing and retrieving data can make a plausible case that it satisfies the OAIS conformance requirements}~\\cite{rosenthal05}, so it is important to be able to reassure yourself, as a data manager, that you have achieved more than simply producing the statement \\qq{we promise not to lose the data}, dressed in OAIS finery.  This is the domain of \\emph{OAIS certification}, and this involves both efforts to define more detailed requirements~\\cite{rosenthal05}, and efforts to devise more stringent and more auditable assessments of an OAIS's actual ability to be appropriately responsive to technology change (see~\\cite{caspar-evaluation09} and \\cite[ch.25]{giaretta11}, and \\prettyref{s:caspar}).  The conjugate of validation is the question of how, as a funder, you reassure yourself that the DMP plan which a project has proposed is actually capable of doing what you (and you hope the project) wish it to do.  Together, these are the domain of \\emph{OAIS auditing}, and this is discussed in \\prettyref{s:oais-audit}  \\subsection{What is \\q{big science}?} \\label{s:bigscience} Big science projects tend to share many features which distinguish them from the way that experimental science has worked in the past.  The differences include big money, big author lists and, most famously, big data: the \\gls{aLIGO} project (for example) will produce of order \\SI1\\PBY, comparable to the \\gls{ATLAS} detector's \\SI{10}\\PBY; the eventual SKA data volumes will dwarf these.  See \\cite[\\S1]{gray10} for extended discussion of the characteristic features of large-scale science.  While the large data volumes bring obvious complications, there are other features of big science which change the way we can approach its data management, and which in fact make the problem easier. \\begin{itemize} \\item Big science projects are often well-resourced, with plenty of relevant and innovative IT experience, engineering management and clear collaboration infrastructure articlulated through \\glspl{MOU}. This means that such projects can develop custom technical designs and implementations, to an extent that would be infeasible for other disciplines. \\item These areas have a long necessary tradition of using shared \\glspl{facility}, so engineering discipline, documented interfaces and SLAs are familiar to the community. \\item Historical experience of \\q{large} data volumes mean everyone knows that ad hoc solutions don't work. Part of the challenge of developing and deploying principled DMP plans in other disciplines is the challenge of persuading funders and senior project members that effective data management is necessary, expensive and technically demanding, and cannot be simply left to junior researchers, however \\q{IT-literate} they may seem to be. This battle is won in disciplines with long experience of large-scale data. \\end{itemize}  In particular, the projects we are focusing on in this project\\dash and what we take the term \\q{big science} to refer to in this report\\dash are the \\q{facilities} and international projects typically funded in the UK by \\gls{STFC}.  A \\q{facility} in this context refers to a (typically large) resource, funded and shared nationally or internationally, which scientists or groups will bid for time on.\\footnote{A telescope's call for proposals is closely analogous to a grant funding call, except that the award will be nights in a forthcoming semester, rather than money.}  The facility will be to some extent a \\q{general purpose} device, such as a telescope or an accelerator like ISIS.  Facilities represent major infrastructural investments, typically enjoy a certain autonomy, and are designed and managed through SLAs. Facilities are generally highly automated, and typically take data directly from the instrument into an archive.  This last point has multiple implications for \\gls{DMP}.  The \\gls{LHC} and \\gls{LIGO} are probably too closely associated with particular goals and collaborations to be naturally termed \\q{facilities}, but they are of the type of international project with the same data challenges.  Our definition of \\q{big science} in this report is, to a first approximation, roughly equivalent to \\q{STFC-funded science}.  STFC is the UK's primary big-science funding council, as it is structured with a particular emphasis on multi-partner collaborative science, less support than the other councils for few-person projects, and budgetary arrangements with the UK Treasury which reflect its exposure to long-term commitments in multiple currencies.  Although most STFC science is \\q{big science} by our definition, the converse is not true, there are examples of such projects funded by both \\gls{EPSRC} and \\gls{NERC}, and we hope that this document will be of use to people in these areas, too.  The most obviously relevant feature of \\q{big science} in our definition is, of course, the \\q{big data} aspect.  Though not a defining feature, it is characteristic of such projects that they are generally willing to deal with data volumes at the upper end of what is feasible, if necessary by designing instruments to produce data volumes no larger than what is predicted to be manageable by the time the instrument finally comes on-line.  Without discounting the technical achievements required by such data rates, the key implication here is that day to day data \\emph{management} is a core concern of the project, which is designed and funded accordingly. There are two key consequences of this, both positive. \\begin{itemize} \\item Data \\emph{preservation}\\dash meaning the continuance of the successfully managed data into the future\\dash is straightforwardly identified as a cousin of the data \\emph{management} problem.  The former problem is not trivial (in a sense expanded on in \\prettyref{s:what-dmp}), since it has distinct goals and, for example, a different budget profile, but some of the more troublesome aspects of \\textit{ab initio} data preservation are handled for free by the necessary existence of a data management infrastructure. \\item In particular, the problems of data ingest, which loom so large in much of the \\gls{DMP} literature, are reduced to the problem of documenting and possibly adjusting archival metadata. \\end{itemize}  Part of the motivation for this present document is the contention that, for technologically sophisticated areas such as this one, the guidance towards the development of a DMP plan can be boiled down to \\qq{Here's a copy of the \\gls{OAIS} spec; get on with it}.  \\subsection{What is \\q{data management and preservation}?} \\label{s:what-dmp} The OAIS specification makes the general remark that \\qq{[t]ransactions among all types of organizations are being conducted using digital forms that are taking the place of more traditional media such as paper.  Preserving information in digital forms is much more difficult than preserving information in forms such as paper and film. This is not only a problem for traditional archives, but also for many organizations that have never thought of themselves as performing an archival function}~\\cite[\\S1.3]{std:oais}.  In the scientific context, \\q{data management} has a somewhat narrower remit: essentially all new scientific data, and a lot of scientific metadata, is \\q{born digital}, and is also born complete, in the sense (expanded in~\\cite[\\S1.7]{gray10}) that the information to be archived is designed and documented in such as way as to support future scientific analysis.  Also\\dash and this is common to most \\glspl{facility} science, and the envy of other disciplines\\dash most large-scale science data is acquired and archived automatically, in a system which must be functioning adequately if the project as a whole is to function at all, so that the matter of data \\emph{preservation} at first appears to be simply a question of copying data from a day-to-day management system into a persistent archive.  But this is not the case.  In large and complicated experiments, the complication of the apparatus makes it hard to communicate into the future a level of understanding sufficient to make plausible use of the data.  This is discussed below, in \\prettyref{s:case-data-preservation}.  This is a useful place to stress that the OAIS definition of the \\gls{Long Term} is simple and pragmatic: the Long Term is, in effect, longer than one technology generation, and thus far enough into the future that the data will have to undergo some storage migration, and that future users will have to depend on documentation rather than human contact with the data creators.  This in turn leads naturally to the observation that data management covers both storage -- the preservation of the bits -- and curation -- the preservation of the knowledge about the bits.  The storage problem is a technical and financial one: we will largely avoid the technical question of which storage technology should be used, save to note that answering this is part of the implementation phase of a DMP plan and that the question must be re-answered by the archive with each technology generation (we discuss storage technology questions very briefly in \\prettyref{s:storage}).  The financial aspect to the storage problem is the question of how much it will cost to store the data into the indefinite future: while storage costs for the few-year short term can be trivially assessed with a couple of hours' work on eBay, the unpredictability of the current long-running decrease in storage prices means that long-term cost estimates are both vital, if a solution is to be sustainable, and very poorly understood.  For a discussion of the estimation of storage costs, see \\prettyref{s:preservation-costs}.  Curation costs, by contrast, are dominated by the front-loaded staff costs for creating Representation Information documentation, and by the non-negligible but broadly predictable staff costs of continuing archive management.    \\ifsnapshot \\begin{wantcomments} Our informal goal in this document is to reassure someone charged with developing a \\gls{DMP} plan that (a) a reasonable framework for approaching the problem already exists in OAIS, that (b) the OAIS recommendation is concrete enough that it is not just waffle, and that (c) X, Y and Z are the things to read to become the local expert, which means that (d) if you're the funder, then Xa, Ya and Za are the questions to ask about the result.  We would be particularly interested in comments which discuss the extent to which we have achieved this goal. \\end{wantcomments} \\fi ",
        "conclusions": ""
    },
    "1208/1208.3562_arXiv.txt": {
        "abstract": "We consider reheating driven by volume modulus decays in the LARGE Volume Scenario. Such reheating always generates non-zero dark radiation through the decays to the axion partner, while the only competitive visible sector decays are Higgs pairs via the Giudice-Masiero term. In the framework of sequestered models where the cosmological moduli problem is absent, the simplest model with a shift-symmetric Higgs sector generates $1.56\\leq \\Delta N_{eff}\\leq 1.74$. For more general cases, the known experimental bounds on $\\Delta N_{eff}$ strongly constrain the parameters and matter content of the models. ",
        "introduction": "The cosmological Standard Model (SM) starts with a period of inflation. During this period, the energy of the universe is dominated by the vacuum energy of a slowly rolling scalar field. At some point inflation ends and, irrespective of the overall particle spectrum or number of hidden sectors, the energy has to be transferred predominantly into thermal relativistic SM degrees of freedom via a process of reheating.  Constraints on this are measured via $N_{eff}$, the effective number of neutrino species. $N_{eff}$ is measured both at BBN and CMB times, and in practice measures the fraction of the total energy density that lies in the thermal photon plasma. At CMB temperatures $N_{eff}$ is determined in terms of the total energy by \\be \\rho_{total} = \\rho_{\\gamma} \\left( 1 + \\frac{7}{8}\\left( \\frac{4}{11} \\right)^{4/3} N_{eff} \\right). \\ee In the SM $N_{eff,BBN} = 3$ and $N_{eff,CMB} = 3.04$, due to partial reheating of the neutrinos from $e^{+} e^{-}$ annihilation. The presence of additional dark radiation, decoupled from the SM and relativistic at both BBN and CMB temperatures, leads to $\\Delta N_{eff} \\equiv N_{eff} - N_{eff,SM} > 0$.  Observations show a mild but consistent preference for $\\Delta N_{eff} > 0$. At CMB times WMAP, ACT and SPT report $N_{eff} = 4.34^{+0.86}_{-0.88},$  $4.56 \\pm 0.75,  3.86 \\pm 0.42$ respectively \\cite{wmapetc}. At BBN times an excess has also been reported but the evidence depends on the relic Helium abundance \\cite{manganoetc}. Based only on D/H, \\cite{NollettHolder} reports $N_{eff} = 3.9 \\pm 0.44$. A recent general overview is \\cite{steigman}.  As the inflationary universe is vacuum energy dominated, dark radiation must arise during or after reheating. In the context of string models of the early universe, there are two main challenges in understanding reheating. The first, the cosmological moduli problem (CMP) \\cite{cmp} is to understand how re-\\emph{heat}-ing can occur at all, and the second is to ensure that it is primarily the SM that is reheated.  We recall first the CMP \\cite{cmp}. String theory contains many moduli associated to the complicated Calabi-Yau geometry. Moduli are typically Planck-coupled scalars which are expected to obtain vevs during inflation, leading to post-inflationary production of moduli through the vacuum misalignment mechanism. Moduli oscillate coherently as matter, redshift slowly and come to dominate the energy density of the universe. As Planck-coupled fields, their characteristic decay rate is $\\Gamma \\sim \\frac{1}{16 \\pi} \\frac{m_{\\phi}^3}{M_P^2}$. A reheating temperature $T \\gtrsim \\mc{O}(1)\\, \\hbox{MeV}$, necessary for BBN, then requires $m_{\\phi} \\gtrsim 30\\, \\hbox{TeV}$. For `generic' models, $m_{\\phi} \\sim m_{3/2} \\sim M_{soft}$, leading to a tension with supersymmetric solutions of the hierarchy problem.  There is a cognate problem with respect to decays to gravitini \\cite{ph0602061}. Even if $m_{\\phi} \\gg M_{soft}$, provided $m_{\\phi} \\gtrsim 2 m_{3/2}$ the decay mode $\\phi \\to \\psi_{3/2} \\psi_{3/2}$ is kinematically open. This decay mode is problematic as for $m_{3/2} \\lesssim 30 \\,\\hbox{TeV}$ the gravitino decays could affect the successful BBN predictions.  The second problem is to ensure that only the SM is reheated. String theory generally contains many extra sectors in addition to the SM. These include additional hidden gauge and matter sectors, as well as light axion-like particles. Excessive branching ratios to these hidden sectors would lead to an overproduction of dark matter or $\\Delta N_{eff} \\gg 1$ and a failure of the BBN predictions. There is also a practical difficulty. Calabi-Yaus have many - easily $\\mc{O}(100)$ - moduli which in generic models of moduli stabilisation have parametrically similar masses and lifetimes. A study of reheating then requires a coupled analysis of all moduli and their decay modes. Such an analysis is not only impractical, it is also highly sensitive to the post-inflationary initial conditions as the relative energy densities in each modulus field depends on the magnitude of the initial modulus misalignment and its non-perturbative production rate at pre-heating. ",
        "conclusions": "The main point to emphasise is that dark radiation is generic and unavoidable in the sequestered LARGE Volume Scenario. This relies only on reheating being driven by decays of the lightest modulus, which always has an open decay mode to its axion partner. The magnitude of this dark radiation depends on assumptions about the visible sector and the number of additional closed string axions, but can easily be at a level consistent with observational hints for $\\Delta N_{eff}$. The bounds on $\\Delta N_{eff}$ can be used to strongly constrain the couplings and matter content of the models, showing that it is very hard to achieve an axiverse in sequestered models. Future observations, such as those expected from PLANCK, will further constrain this scenario.  \\emph{Note added: This paper is submitted simultaneously to the related work \\cite{TakahashiHigaki}.} \\subsection*"
    },
    "1208/1208.3048_arXiv.txt": {
        "abstract": "{The Hundred-Thousand-Proper-Motion (HTPM) project will determine the proper motions of $\\sim$$113,500$ stars using a $\\sim$23-year baseline. The proper motions will be based on space-based measurements exclusively, with the Hipparcos data, with epoch $1991.25$, as first epoch and with the first intermediate-release Gaia astrometry, with epoch $\\sim$$2014.5$, as second epoch. The expected HTPM proper-motion standard errors are $30$--$190$~$\\mu$as~yr$^{-1}$, depending on stellar magnitude.} {Depending on the astrometric characteristics of an object, in particular its distance and velocity, its radial velocity can have a significant impact on the determination of its proper motion. The impact of this perspective acceleration is largest for fast-moving, nearby stars. Our goal is to determine, for each star in the Hipparcos catalogue, the radial-velocity standard error that is required to guarantee a negligible contribution of perspective acceleration to the HTPM proper-motion precision.} {We employ two evaluation criteria, both based on Monte-Carlo simulations, with which we determine which stars need to be spectroscopically (re-)measured. Both criteria take the Hipparcos measurement errors into account. The first criterion, the Gaussian criterion, is applicable to nearby stars. For distant stars, this criterion works but returns overly pessimistic results. We therefore use a second criterion, the robust criterion, which is equivalent to the Gaussian criterion for nearby stars but avoids biases for distant stars and/or objects without literature radial velocity. The robust criterion is hence our prefered choice for all stars, regardless of distance.} {For each star in the Hipparcos catalogue, we determine the confidence level with which the available radial velocity and its standard error, taken from the XHIP compilation catalogue, are acceptable. We find that for \\color{black} $97$ \\color{black} stars, the radial velocities available in the literature are insufficiently precise for a $68.27\\%$ confidence level. If requiring this level to be $95.45\\%$, or even $99.73\\%$, the number of stars increases to \\color{black} $247$ \\color{black} or \\color{black} $382$\\color{black}, respectively. We also identify \\color{black} $109$ \\color{black} stars for which radial velocities are currently unknown yet need to be acquired to meet the $68.27\\%$ confidence level. For higher confidence levels ($95.45\\%$ or $99.73\\%$), the number of such stars increases to \\color{black} $1,071$ \\color{black} or \\color{black} $6,180$\\color{black}, respectively.} {To satisfy the radial-velocity requirements coming from our study will be a daunting task consuming a significant amount of spectroscopic telescope time. The required radial-velocity measurement precisions vary from source to source. Typically, they are modest, below $25$~km~s$^{-1}$, but they can be as stringent as $0.04$~km~s$^{-1}$ for individual objects like Barnard's star. Fortunately, the follow-up spectroscopy is not time-critical since the HTPM proper motions can be corrected {\\it {a posteriori}} {\\rm{once (improved) radial velocities become available.}}} ",
        "introduction": "Gaia \\citep[e.g.,][]{2001A&A...369..339P,2008IAUS..248..217L} is the upcoming astrometry mission of the European Space Agency (ESA), following up on the success of the Hipparcos mission \\citep{1997ESASP1200.....P,1997A&A...323L..49P,2009aaat.book.....P}. Gaia's science objective is to unravel the kinematical, dynamical, and chemical structure and evolution of our galaxy, the Milky Way \\citep[e.g.,][]{2010MNRAS.408..935G}. In addition, Gaia's data will revolutionise many other areas of (astro)physics, e.g., stellar structure and evolution, stellar variability, double and multiple stars, solar-system bodies, fundamental physics, and exo-planets \\citep[e.g.,][]{2008IAUS..248...59P,2008P&SS...56.1812T,2010IAUS..261..306M,2011EAS....45..161E,2011EAS....45..273S,2011PhRvD..84l2001M}. During its five-year lifetime, Gaia will survey the full sky and repeatedly observe the brightest 1,000 million objects, down to $20^{\\rm th}$ magnitude \\citep[e.g.,][]{2010SPIE.7731E..35D}. Gaia's science data comprises absolute astrometry, broad-band photometry, and low-resolution spectro-photometry. Medium-resolution spectroscopic data will be obtained for the brightest 150 million sources, down to $17^{\\rm th}$ magnitude. The final Gaia catalogue, due in $\\sim$$2021$, will contain astrometry (positions, parallaxes, and proper motions) with standard errors less than $10$~micro-arcsecond ($\\mu$as, $\\mu$as~yr$^{-1}$ for proper motions) for stars brighter than $12^{\\rm th}$ magnitude, $25~\\mu$as for stars at $15^{\\rm th}$ magnitude, and $300~\\mu$as at magnitude 20 \\citep{2012Ap&SS.tmp...68D}. Milli-magnitude-precision photometry \\citep{2010A&A...523A..48J} allows to get a handle on effective temperature, surface gravity, metallicity, and reddening of all stars \\citep{2010MNRAS.403...96B}. The spectroscopic data allows the determination of radial velocities with errors of $1~{\\rm km~s}^{-1}$ at the bright end and $15~{\\rm km~s}^{-1}$ at magnitude 17 \\citep{2005MNRAS.359.1306W,2011EAS....45..189K} as well as astrophysical diagnostics such as effective temperature and metallicity for the brightest few million objects \\citep{2011A&A...535A.106K}. Clearly, these performances will only be reached with a total of five years of collected data and only after careful calibration.  Intermediate releases of the data -- obviously with lower quality and/or reduced contents compared to the final catalogue -- are planned, the first one around two years after launch, which is currently foreseen for the second half of 2013. The Hundred-Thousand-Proper-Motion (HTPM) project \\citep{FM-040}, conceived and led by Fran\\c{c}ois Mignard at the Observatoire de la C\\^{o}te d'Azur, is part of the first intermediate release. Its goal is to determine the absolute proper motions of the $\\sim$$113,500$ brightest stars in the sky using Hipparcos astrometry for the first epoch and early Gaia astrometry for the second. Clearly, the HTPM catalogue will have a limited lifetime since it will be superseded by the final Gaia catalogue in $\\sim$$2021$. Nevertheless, the HTPM is a scientifically interesting as well as unique catalogue: the $\\sim$23-year temporal baseline, with a mean Hipparcos epoch of $1991.25$ and a mean Gaia epoch around $2014.5$, allows a significant improvement of the Hipparcos proper motions, which have typical precisions at the level of $1$~milli-arcsecond~yr$^{-1}$ (mas~yr$^{-1}$): the expected HTPM proper-motion standard errors\\footnote {The HTPM proper motions will be limited in precision by the Hipparcos parallax uncertainties, which are typically $\\sim$1~mas (the typical HTPM proper-motion standard error is hence $1~{\\rm mas} / 23~{\\rm yr} \\approx 40~\\mu{\\rm as~yr}^{-1}$). The first intermediate-release Gaia catalogue is based on just $\\sim$$12$ months of data, which is generally insufficient to unambiguously lift the degeneracy between proper motion and parallax for all stars. The underlying astrometric global iterative solution \\citep{2012A&A...538A..78L} will hence be based on a two- rather than five-parameter source model, fitting for position $(\\alpha, \\delta)$ at the mean Gaia epoch only. The Hipparcos parallax is hence needed to correct the Gaia transit observations for the parallactic effect allowing to transform apparent directions into barycentric positions.} are $40$--$190~\\mu$as~yr$^{-1}$ for the proper motion in right ascension $\\mu_{\\alpha^*}$ and $30$--$150$~$\\mu$as~yr$^{-1}$ for the proper motion in declination $\\mu_{\\delta}$, primarily depending on magnitude (we use the common Hipparcos notation $\\alpha^* = \\alpha \\cdot \\cos\\delta$; \\citealp[Section 1.2.5]{1997ESASP1200.....P}). A clear advantage of combining astrometric data from the Hipparcos and Gaia missions is that the associated proper motions will be, by construction and IAU resolution, in the system of the International Celestial Reference System (ICRS), i.e., the proper motions will be absolute rather than relative. In this light, it is important to realise that massive, modern-day proper-motion catalogues, such as UCAC-3 \\citep{2010AJ....139.2184Z}, often contain relative proper motions only and that they can suffer from substantial, regional, systematic distortions in their proper-motion systems, up to levels of 10~mas~yr$^{-1}$ or more \\citep[e.g.,][]{2008A&A...488..401R,2010AJ....139.2440R,2011RAA....11.1074L}.  It is a well-known geometrical feature, for instance already described by \\citeauthor{1900AN....154...65S} in \\citeyear{1900AN....154...65S}, that for fast-moving, nearby stars, it is essential to know the radial velocity for a precise measurement and determination of proper motion. In fact, this so-called secular or perspective acceleration on the sky was taken into account in the determination of the Hipparcos proper motions for 21 stars \\citep[Section 1.2.8]{1997ESASP1200.....P} and the same will be done for Gaia, albeit for a larger sample of nearby stars. Clearly, the inverse relationship also holds: with a precise proper motion available, a so-called astrometric radial velocity can be determined, independent of the spectroscopically measured quantity (see \\citealt{2003A&A...401.1185L} for a precise definition and meaning of [astrometric] radial velocity). With this method, \\citet{1999A&A...348.1040D} determined\\footnote {These authors also describe two other methods to derive astrometric radial velocities, namely by measuring changing annual parallax or by measuring changing angular extent of a moving group of stars \\citep{2002A&A...381..446M}. The latter method also provides, as a bonus, improved parallaxes to moving-group members \\citep[e.g.,][]{1999MNRAS.310..585D,2001A&A...367..111D}.} the astrometric radial velocities for 17 stars, from Hipparcos proper motions combined with Astrographic Catalogue positions at earlier epochs. Although \\citet{1999A&A...348.1040D} reached relatively modest astrometric-radial-velocity precisions, typically a few tens of km~s$^{-1}$, their results are interesting since they provide direct and independent constraints on various physical phenomena affecting spectroscopic radial velocities, for instance gravitational redshifts, stellar rotation, convection, and pulsation. In our study, however, we approach (astrometric) radial velocities from the other direction since our interest is to determine accurate HTPM proper motions which are not biased by unmodeled perspective effects. In other words: we aim to establish for which stars in the forthcoming HTPM catalogue the currently available (spectroscopic) radial velocity and associated standard error are sufficient to guarantee, with a certain confidence level, a negligible perspective-acceleration-induced error in the HTPM proper motion. For stars without a literature value of the radial velocity, we establish whether -- and, if yes, with what standard error -- a radial velocity needs to be acquired prior to the construction of the HTPM catalogue. Section~\\ref{sec:XHIP} describes the available astrometric and spectroscopic data. The propagation model of star positions is outlined in Section~\\ref{sec:propagation_model}. We investigate the influence of the radial velocity on HTPM proper motions in Section~\\ref{sec:rv_influence} and develop two evaluation criteria in Section~\\ref{sec:evaluation_criteria}. We employ these in Section~\\ref{sec:application}. We discuss our results in Section~\\ref{sec:discussion} and give our final conclusions in Section~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have conducted a study of the requirements for the availability of radial velocities for the Hundred-Thousand-Proper-Motion (HTPM) project \\citep{FM-040}. This unique project will combine Hipparcos astrometry from $1991.25$ with early-release Gaia astrometry ($\\sim$$2014.5$) to derive long-time-baseline and hence precise proper motions. For the nearest, fast-moving stars, the perspective acceleration of the objects on the sky requires the presence of radial velocities for the derivation of the proper motions. We have quantitatively determined, for each star in the Hipparcos catalogue, the precision of the radial velocity that is required to ensure that the perspective-acceleration-induced error in the HTPM proper motion caused by the radial-velocity error is negligible. Our method takes the Hipparcos measurement errors into account and allows the user to specify his/her own prefered confidence level (e.g., $68.27\\%$, $95.45\\%$, or $99.73\\%$). The results are available in Table~\\ref{tab:datafile} (Appendix~\\ref{sec:datafile}), which is published electronically only. We have compared the radial-velocity-precision requirements to the set of $46,392$ radial velocities contained in the XHIP compilation catalogue \\citep{2012AstL...38..331A} and find that, depending on the confidence level one wants to achieve, hundreds to thousands of stars require spectroscopic follow-up. The highest-priority targets are \\color{black} $206$ \\color{black} objects with a confidence level below $68.27\\%$; \\color{black} $97$ \\color{black} of them have a known but insufficiently precise radial velocity while the remaining \\color{black} $109$ \\color{black} objects have no literature radial velocity in the XHIP compilation catalogue at all. The typical brightness of the objects requiring their radial velocity to be (re-)determined is $H\\!p \\approx 8$--$12$~mag and the radial-velocity precisions vary drastically, ranging from $0.04$~km~s$^{-1}$ for the most extreme case (HIP$87937$, also known as Barnard's star) to a few tens of km~s$^{-1}$. With only few exceptions, the spectral types are K and M; \\color{black} $73$\\% \\color{black} of them are in the south. Gaia's Radial-Velocity Spectrometer (RVS; \\citealp{2011EAS....45..181C}) will deliver radial velocities for all stars in the HTPM catalogue with Gaia-end-of-mission precisions below a few km~s~$^{-1}$ (and $\\sim$$10$~km~s~$^{-1}$ for early-type stars; \\citealp{2012Ap&SS.tmp...68D}); however, these performances require full calibration of the instrument and data and hence will most likely only be reached in the final Gaia data release, at which time the HTPM proper motions will be superseded by the Gaia proper motions. Fortunately, the spectroscopic follow-up is not time-critical in the sense that the HTPM catalogue will be published with information (sensitivity coefficients and reference parallax and radial velocity) to correct the proper motions {\\it a posteriori} when (improved) radial velocities become available. We finally note that the spectroscopic follow-up requirements for the HTPM proper motions quantified in this work will be dwarfed by the requirements coming from the end-of-mission Gaia proper motions, to be released around $\\sim$$2021$: for instance for the stars in the HTPM catalogue, for which the HTPM proper-motion \\color{black} standard \\color{black} errors are $30$--$190$~$\\mu$as~yr$^{-1}$, the Gaia proper-motion standard errors reach the bright-star floor around $3$--$4~\\mu$as~yr$^{-1}$ \\citep{2012Ap&SS.tmp...68D}, which means that the spectroscopic requirements for the correction of perspective acceleration in the Gaia astrometry will be a factor $\\sim$$10$--$50$ more demanding."
    },
    "1208/1208.0360_arXiv.txt": {
        "abstract": "{ X-ray reflection off the accretion disc surrounding a black hole, together with the associated broad iron K$\\alpha$ line, has been widely used to constrain the innermost accretion-flow geometry and black hole spin. Some recent measurements have revealed steep reflection emissivity profiles in a number of active galactic nuclei and X-ray binaries. } { We explore the physically motivated conditions that give rise to the observed steep disc-reflection emissivity profiles. } { We perform a set of simulations based on the configuration of a possible future high-resolution X-ray mission. Computations are carried out for typical X-ray bright Seyfert-1 galaxies. } { We find that steep emissivity profiles with $q\\sim 4-5$ (where the emissivity is $\\epsilon (r) \\propto r^{-q}$) are produced considering either i) a lamp--post scenario where a primary compact X--ray source is located close to the black hole, or ii) the radial dependence of the disc ionisation state. If both effects are taken into account, emissivity profiles as steep as $q \\sim 7$ can be obtained from X--ray spectra modelled via conventional reflection models. We also highlight the role of the reflection angular emissivity: the radial emissivity index $q$ is overestimated when the standard limb--darkening law is used to describe the data. } { Very steep emissivity profiles with $q \\geq 7$ are naturally obtained by applying reflection models % that take into account radial profile $\\xi (r)$ of the disc ionisation induced by a compact X--ray source located close to the central black hole. } ",
        "introduction": "%  The innermost black-hole accretion discs of active galactic nuclei (AGNs) and black hole binaries may be revealed through their X-ray radiation released as either the thermal radiation of the disc \\citep{1997ApJ...482L.155Z, 2006csxs.book..157M} or the result of the inverse Compton scattering of the thermal photons on the relativistically moving electrons in the hot corona above the disc \\citep{1976ApJ...204..187S, 1991ApJ...380L..51H}. While the first is relevant only to stellar-mass black holes where the accretion disc is heated to very high temperatures (around $10^7$\\,K), the second is common for any accreting black hole over the entire possible mass scale.  Some fraction of the photons scattered in the corona reflect off the disc surface before reaching the observer. This so-called reflection spectrum is characterised by a Compton hump at energies of around $20-40$\\,keV, fluorescent lines of which the iron K$\\alpha$ line at $6.4-6.97$\\,keV is the most prominent, and a soft excess below 1\\,keV. The overall spectral shape depends on the ionisation of the disc surface. The spectrum is then smeared by the relativistic effects including energy shift, aberration, and light-bending. The inner disc radius, inclination angle, average ionisation state, and reflection-emissivity radial profile can in principle be determined from the resulting X--ray spectral shape. The inner disc radius plays a particularly important role. As shown e.g. by \\citet{2008ApJ...675.1048R}, the inner-disc reflection radius can be associated with the innermost stable circular orbit, which only depends on the black hole spin. Hence, X--ray reflection spectra can be used to estimate the black hole spin in both AGN and black hole binaries \\citep[see e.g.][and references therein]{2003PhR...377..389R, 2009ApJ...697..900M, 2009ApJ...702.1367B}.  The black hole spin measurements are thus influenced by the geometry of the disc-illuminating corona and the local properties of the disc that affect the re-processing and re-emission of the incident photons. Current relativistic kernels that are applied to reflection (and/or iron line) models to include the relativistic effects on the spectral shape \\citep{1991ApJ...376...90L, 2004ApJS..153..205D, 2004MNRAS.352..353B, 2006ApJ...652.1028B} are based on a series of simplified assumptions and, in particular, they assume a single/broken power-law form of the radial reflection emissivity and an angular emissivity law defined by a simple analytical formula, most frequently employing a limb-darkening profile \\citep{1991ApJ...376...90L}.   In \\citet{2009A&A...507....1S}, we pointed out that the choice of a particular angular emissivity law can affect the derived relativistic parameters, including the black hole spin. In the present work, we extend our analysis to the radial emissivity law and, in particular, we investigate the physical conditions under which steep radial emissivity profiles are produced by considering the effects of (i) primary X-ray source location, (ii) radial disc ionisation profile, and (iii) angular emissivity law.   The intrinsic disc emissivity is naturally expected to decrease with increasing distance, i.e. the reflection emissivity is \\begin{equation} \\epsilon(r) = r^{-q}, \\end{equation} where $q$ is the emissivity index that can be constant over all radii or a varying quantity. The thermal dissipation of the disc decreases as $r^{-3}$ \\citep{1973A&A....24..337S, 1973blho.conf..343N}. Therefore, the simplest assumption is postulating the same dependence for the reflection. The more energetic photons are injected into the innermost regions, and so, more intense irradiation of the disc occurs there. In addition, assuming a point-like X-ray source at height $h$ on the disc axis, the irradiation of the disc in the absence of any relativistic effect is proportional to $(r^2 + h^2)^{-3/2}\\propto r^{-3}$, as shown e.g. by \\citet{1997ApJ...488..109R}. An emissivity profile with $q=3$ is therefore considered as {\\it standard}, while steeper/flatter indices may need to be explained.   Steep emissivity profiles have been reported from the analysis of the X-ray spectra of other AGN, such as MCG\\,-6-30-15 \\citep{2002MNRAS.335L...1F, 2004MNRAS.348.1415V, 2007PASJ...59S.315M}, 1H0707-495 \\citep{2009Natur.459..540F, 2010MNRAS.401.2419Z, 2011MNRAS.414.1269W, 2012MNRAS.422.1914D}, and IRAS\\,13224-3809 \\citep{2010MNRAS.406.2591P}, as well as black hole binaries, such as XTE\\,J1650-500 \\citep{2004MNRAS.351..466M}, GX\\,339-4 \\citep{2007ARA&A..45..441M}, and Cyg\\,X-1 \\citep{2012arXiv1204.5854F}. The measured indices reach values up to $q \\approx 7$.  To provide a physical picture of the steep radial emissivity in MCG\\,-6-30-15, \\citet{2001MNRAS.328L..27W} invoke strong magnetic stresses that should act in the innermost region of the system. This should correspond to the enhanced dissipation of a considerable amount of energy in the accretion disc at small radii. If the magnetic field lines thread the black hole horizon, the dissipation could be triggered by the magnetic extraction of the black hole rotational energy, perhaps via the Blandford-Znajek effect \\citep{1977MNRAS.179..433B}, but it could be also supplemented by a rather efficient slowing of the rotation, as also seen in recent GRMHD simulations \\citep[e.g.][]{2010MNRAS.408..752P}. The efficiency of the competing processes still needs to be assessed.  \\citet{2000MNRAS.312..817M} % examined whether the required steep emissivity law as well as the predicted equivalent width of the cold reflection line of iron and the Compton reflection component can be reproduced in a phenomenological (lamp-post) model where the X-ray illuminating source is located on the common symmetry axis of the black hole and the equatorial accretion disc. They suggested that the radial emissivity function of the reflection component steepens when the height parameter of the primary irradiation source decreases. The enhanced anisotropy of the primary X-rays was identified as a likely agent acting in this process. The emissivity in the XMM-Newton spectrum of MCG\\,-6-30-15 was successfully reproduced by adopting this lamp-post geometry \\citep{2002A&A...383L..23M, 2003MNRAS.344L..22M, 2008MNRAS.386..759N}.    To explain the steep radial emissivity, we explore several simple test models and analyse them with the simulated data. The paper is organised as follows. Sect.~\\ref{data} describes the data generation. The effects of the lamp-post geometry of the corona, the angular directionality, and the radially stratified ionisation on the radial emissivity are discussed in Sects.~\\ref{section_lamp_post}, ~\\ref{section_angular}, and ~\\ref{section_radion}, respectively. An example of the combined effect is presented in Sect.~\\ref{section_combi}. The obtained results are discussed in Sect.~\\ref{discussion} and summarised in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{discussion}%   We have addressed the steep radial emissivities that have been detected in the reflection components of the X-ray spectra of active galaxies and black-hole binaries. We investigated some possible explanations, and to this end performed several simulations to reveal the degeneracies of the radial emissivity with other parameters and the intrinsic assumptions of the relativistic reflection model. We realised that the steep radial emissivity may be explained by either (i) the geometrical properties of the disc-illuminating corona, (ii) the use of an improper model assumption about the angular directionality, or (iii) radially structured ionisation. The first puts rather extreme requirements on the corona. It needs to be very bright and occur at~a very low height above the black hole. The second is likely due to use of an improper prescription for the emission directionality in the black hole accretion disc. The last one is related to a probable radial dependence of the disc ionisation, which plays a significant role in the total shape of the reflection spectrum.  \\subsection{Lamp-post scenario}  The steep radial emissivity may be related to the properties of the disc-illuminating corona as suggested before by \\citet{2001MNRAS.328L..27W}. The geometry of the emitting region certainly plays a significant role. A very centrally localised source at a low height above the black hole horizon would irradiate the disc mainly in its central region. The illumination in this area is greatly enhanced by the gravitational light-bending effect \\citep[Dov\\v{c}iak et al., in prep.]{2004MNRAS.349.1435M, 2008MNRAS.386..759N, 2011MNRAS.414.1269W}.  To achieve steep radial emissivity, which is assumed to be proportional to the illumination, the source must be sufficiently close to the black hole. However, in this case the primary emission has to be extremely bright because only a small fraction would overcome the strong gravitational pull of the black hole and reach the observer \\citep[see Fig.\\,2 in][]{2011ApJ...731...75D}. The importance of these effects declines sharply with height. At heights $h \\gtrsim 3\\,r_{\\rm g}$, the radial emissivity profile is similar to the simple power-law with the standard value ($q=3$). For even larger heights, the irradiation profile is more complicated \\citep[see Fig.\\,3 and 4 in][]{2011ApJ...731...75D}. It decreases steeply only very close to the black-hole horizon, then becomes rather flat ($q<3$) when still in the inner parts of the disc and finally reaches the standard value far from the centre.     \\subsection{Angular directionality}  For the angular emissivity, the limb darkening law is frequently used. Several simulations, however, suggest that the directionality is the opposite of limb darkening \\citep[see e.g.][ and references therein]{2011A&A...527A..47R}. The emission angle in the innermost region of the disc is always very high owing to strong aberration. The flux contribution from this region is therefore underestimated by models with limb darkening if the angular emissivity is indeed different. This effect could lead to an approximately $20\\,\\%$ overestimation of either the spin or the inner radial emissivity parameter. \\citet{2009A&A...507....1S} reanalysed the XMM-Newton observation of MCG\\,-6-30-15 and showed that the radial emissivity might be a more sensitive parameter to the angular directionality than the spin. This is especially true when the spin value itself is very high (close to one).   \\begin{figure}% \\includegraphics[width=0.5\\textwidth]{radion_constrho_lph015_fixp_LD_contaq_2.eps} \\caption{Contour plot of the radial emissivity parameter $q$ and the spin $a$. The data were created by a lamp-post model with the height $h=1.5\\,r_{\\rm g}$, radially stratified ionisation, and isotropic angular emissivity. A single-ionisation model with limb darkening and power-law description for the radial emissivity was used to fit the data. The best-fit parameters are indicated by a small cross within the contours.} \\label{combi_contaq} \\end{figure}   \\subsection{Radially structured ionisation}  We also discussed the impact of the probable radial dependence of the disc surface ionisation. The disc illumination by a corona is commonly assumed to be stronger in the innermost regions. We therefore assumed that the ionisation is higher in the innermost region and decreases with radius. We did not consider other aspects that affect the ionisation structure of the disc, such as the density profile, vertical structure, and thermal processes (the last one being especially relevant for the stellar-mass black hole binaries). We simply assumed that the radial dependence of the irradiation is the dominating effect in determining the ionisation of the disc surface. This is especially true when the lamp-post scenario for a source of low height is considered.  We note, however, that the density may still play an important role around the marginally stable orbit. While the Novikov-Thorne profile of the density diverges there \\citep{1973blho.conf..343N}, i.e. the ionisation would decrease to zero, the accretion disc may extend below the marginally stable orbit owing to the presence of a magnetic field, as in the GRMHD simulations of \\citet{2006ApJ...641..103H}. The density does not increase to an infinite value, but instead decreases close to the marginally stable orbit \\citep[see also][]{2008ApJ...675.1048R}. This region  was, for clarity, not shown in Figure~\\ref{xi_profile} but we plan to investigate it in more detail in a follow-up paper.   By fitting the simulated data, we realised that the radial decrease in the disc ionisation may account for the radial-emissivity steepness equally as well as the assumption of the centrally localised corona. For the case of an isotropically illuminated disc, we obtained radial emissivities of $q \\approx 4-5$, i.e. values similar to those for the lamp-post irradiation of a cold disc with the height $h=1.5 r_{\\rm g}$. This is due to the different shapes of the reflection model for different ionisation parameters. The softer, more ionised reflection comes from the innermost part of the disc. When a simplified model with a single ionisation is used to fit the data, it may lead to a significant underestimation of the flux from the innermost regions, which is then artificially compensated for by a steep value of the radial emissivity profile in the model.  The main difference between the single-ionisation model and the data created by a complex ionisation occurs at the iron-line energy band and its edge (see Figs.~\\ref{iso_plld_step} and \\ref{combi_plld_step}). While the shape of the ``complex'' reflection continuum may be mimicked by a~modest value of the ionisation and steep radial emissivity, the iron line peak due to reflection from the cold distant parts of the disc is not incorporated in the simplified model. An additional emission line is then needed to model the residuals. In our test case, the equivalent width of such a line was found to be about $40$\\,eV. Several AGN spectra, such as MCG\\,-6-30-15 \\citep{2002MNRAS.335L...1F}, have been found to contain, together with the broad disc component, a narrow iron line associated with reflection from a distant torus. A typical equivalent width is $\\sim 100$\\,eV. The reflection from the outer parts of the accretion disc might provide a significant contribution to this line when the single-ionisation reflection model is applied to the disc-line modelling.   In our simulations, we considered the radial dependence of the ionisation. However, a similar effect on the emissivity profile may also be caused by a vertically structured accretion disc where the core is cold but the outer layers are hot, i.e. strongly ionised \\citep{2000ApJ...537..833N}. \\citet{2001ApJ...546..406N} showed that in the case of AGNs, when the illuminating flux is much higher than the thermal radiation of the disc, the hot skin of the disc is completely ionised and most of the re-processing occurs in the disc core. However, \\citet{2012MNRAS.422.1914D} revealed the significant presence of highly ionised reflection together with relatively cold reflection from the disc in the X-ray spectrum of a narrow-line Seyfert 1 galaxy 1H0707-495. The double-reflection model is shown in their Figure~6. Owing to their different shapes, the ionised reflection contributes more to the red wing of the observed iron line than the cold component. Neglecting the ionised component would therefore result in the measurement of a steeper radial emissivity index.  Finally, we note that the simulations with the radially changing ionisation for the isotropic irradiation were done with the assumption of a ``standard'' value for the radial emissivity, $q=3$. However, non-thermal coronal emission does not necessarily need to behave in the same way as the thermal dissipation of the disc. The interaction between the disc and the corona is more complicated, including the radiation and magnetic processes \\citep[see e.g.][]{1991ApJ...380L..51H, 2004A&A...428..353C, 2006AN....327..977G, 2011A&A...527A..47R}. In particular, when the magnetic field is considered, the intrinsic profile might already be as steep as $r^{-5}$ \\citep{2005ApJ...635..167K}. Further investigation of these poorly understood disc-corona interactions is therefore desirable.     \\subsection{The combined effect}  All the effects that we have discussed are not independent of each other. Irradiation by a compact centrally localised source is very anisotropic \\citep{2004MNRAS.349.1435M}. The central regions are more illuminated and thus, more likely to be ionised than the outer parts of the accretion disc. The effect of a hypothetical radial stratification of the disc ionisation is stronger than in the case of isotropic irradiation. We therefore performed a final simulation including both effects together. We simulated the data with the lamp-post geometry and the ionisation stratification calculated from the theoretical irradiation profile (assuming constant density). We fitted the generated spectrum with a single-ionisation model with a broken power-law for radial emissivity. As expected, we obtained significantly steeper radial emissivity, $q \\approx 7$ (see the right panel of Table~\\ref{table_model} and Figure~\\ref{combi_contaq}).  Angular directionality plays an important role in the total emissivity. We showed that the model with limb darkening overestimates the index of the radial emissivity profile by $5-10\\%$ (see the right panel of Figure~\\ref{combi_plld_step}). The effect might be even stronger because our numerical calculations of local emissivity suggest that there is a limb-brightening effect, while we conservatively used isotropic directionality in our simulations. Moreover, in the case of anisotropic illumination by a lamp-post source, the emissivity also depends on the incident angle with a similar limb-brightening profile (Dov\\v{c}iak et al., in prep.). The highest output of the reflected emission is thus obtained when both angles are grazing (perpendicular to the disc normal). This happens especially very close to the black hole owing to the large aberration that is caused by the extreme velocities of the matter in the inner disc. Hence, the flux from the innermost region might be significantly enhanced and thus the radial-emissivity index overestimated."
    },
    "1208/1208.2153_arXiv.txt": {
        "abstract": "The only prominent line of singly ionized helium in the visible spectral range, He\\,{\\sc ii}\\,4686\\AA, is observed together with the He\\,{\\sc i}\\,5015\\AA{} singlet and the He\\,{\\sc i}\\,4471\\AA {} triplet line in solar prominences. The Na\\,D$_2$ emission is used as a tracer for He\\,{\\sc ii} emissions which are sufficiently bright to exceed the noise level near $10^{-6}$ of the disk-center intensity. The so selected prominences are characterized by small non-thermal line broadening and almost absent velocity shifts, yielding narrow line profiles without wiggles. The reduced widths [$\\Delta\\lambda_D/\\lambda$] of He\\,{\\sc ii}\\,4686\\AA{} are 1.5 times broader than those of He\\,{\\sc i}\\,4471\\AA{} triplet and 1.65 times broader than those of He\\,{\\sc i}\\,5015\\AA{} singlet. This indicates that the He lines originate in a prominence--corona transition region with outwards increasing temperature. ",
        "introduction": "%  As the second most abundant element, helium plays an essential role in astrophysics. Nevertheless, its spectrum is poorly understood. The faint He\\,{\\sc ii}\\,4685.7\\AA{} line is of particular interest, since it is the only important He\\,{\\sc ii} line that can be observed with ground-based telescopes allowing higher spectral resolution than currently achieved for EUV He\\,{\\sc ii} lines from space: \\citet{stellmacher03} find that the SUMER spectrograph (designed for broad coronal lines) does not resolve the narrow emissions from cool prominences, even after application of a maximum instrumental profile.  He\\,{\\sc ii}\\,4685.7\\AA{} has been observed in prominences during eclipses \\citep{sotirovski65,poletto67}. Line-profile analyses, based on moderately resolved spectra from coronographs \\citep{hirayama72,hirayama74} indicate that He\\,{\\sc ii}\\,4685.7\\AA{} originates from the same (cool) prominence regions as the usually observed Balmer, He\\,{\\sc i} and metallic lines. \\citet{zirin59} observe in ''flare-like loop'' ({\\it i.e.} highly active and ''hot'') prominences the width of He\\,{\\sc ii}\\,4685.7\\AA{} larger than that of He\\,{\\sc i}\\,4471.5\\,\\AA, however, they neither discuss quiescent (''cool'') prominences nor the atomic fine-structure broadening of He\\,{\\sc ii} which may explain most of that excess. A detailed analysis of the He\\,{\\sc i} and He\\,{\\sc ii} lines in solar prominences requires high spectral resolution, high signal-to-noise ratio and careful absolute calibration, which is difficult to achieve \\citep[see][]{illing75}, but is possible with modern CCD techniques.  The high ionization and excitation energy of 25 and 48 eV suggests that He\\,{\\sc ii}\\,4685.7\\AA{} may preferentially occur in hot prominences ($T_{\\rm kin}>8000$K) which are observed to have a high He-to-Balmer emission ratio and to be highly structured \\citep{stellmacher94}. These, however, usually show important velocity shifts which disperse the line profiles. We therefore suppose that the chance to measure He\\,{\\sc ii}\\,4685.7\\AA{} increases {\\it if the emitted photons are concentrated in wavelength, i.e. yield narrow line profiles free from spatial Doppler shifts}. Such emissions should occur in prominences with negligible macro-velocities and low non-thermal line broadening, which are known to be bright in H${\\alpha}$ though with a low He-to-Balmer ratio \\citep{engvold71,stellmacher95}. Such prominences show significant Na\\,D and Mg\\,b emission and saturated H${\\alpha}$ (and even H${\\beta}$) profiles but did so far not allow a He\\,{\\sc ii}\\,4685.7\\AA{} line profile analysis \\citep{stellmacher05}. ",
        "conclusions": "\\subsection{Line Radiance}  In order to correctly describe the helium spectrum in the prominence plasma, one has to consider the statistical equilibrium of ortho-, para- and ionized helium and their interaction ({\\it e.g.} \\opencite{labrosse10}). Observations of spectrally well resolved line profiles are required for the understanding of the helium spectrum. The optically thin He\\,{\\sc ii}\\,4685.7\\AA{} line is poorly documented from observations. Its excitation is sensitive to radiation and collisions \\citep{yakovkin71,labrosse10}, but also the influence of turbulence and of flows may enhance the He\\,{\\sc ii} emission \\citep{jordan97,patsourakos02}. A blending with Ni\\,{\\sc i}\\,4686.22\\AA{}, mentioned by \\citet{worden73} in their study of chromospheric emissions, is not indicated in our prominence observations.  \\begin{table} \\caption{Total line emission (radiance) [erg~s$^{-1}$~cm$^{-2}$~sterad$^{-1}$] in comparison with eclipse data from \\citet{sotirovski65} and from \\citet{poletto67}.} \\begin{tabular}{lcccccccc}                       % \\hline position & W$22^\\circ$N & E$14^\\circ$N & E$30^\\circ$N & E$31^\\circ$N & E$23^\\circ$N & Soti- & Poletto \\\\ obs.date       &2\\,Aug.&4\\,Aug.&9\\,Aug.\\,a&9\\,Aug.\\,b&13\\,Aug.& rovski& + Rigutti\\\\ \\hline \\\\ He\\,{\\sc ii}\\,4685.7         & 105  &  35  &   38  &   --  &   13  &  4.0  &  5.5  \\\\ He\\,{\\sc i}\\,4471.5\\,(tr)    & 5119 & 1440 &  1940 &  1300 &  745  &  105  &  80   \\\\ He\\,{\\sc i}\\,5015.7\\,(si)    &  --  &  --  &  198  &   180 &   82  &  19   &  --   \\\\ Fe\\,{\\sc ii}\\,5018.4         &  --  &  --  &  225  &   340 &   37  &   --  &  --   \\\\ Ti\\,{\\sc ii}\\,4468.4         &  85  &  185 &   64  &   105 &   31  &   --  &  --   \\\\ Na\\,D$_2$\\,5890              & 1010 &  540 &  715  &   970 &   70  &   --  &  --   \\\\ \\\\  tripl/He\\,{\\sc ii}           &   49 &   41 &   50  &   --  &   63  &  26   &   15  \\\\ tripl/singl                  &  --  &  --  &   9.8 &  7.2  &  9.1  &  5.7  &  --   \\\\ tripl/Na\\,D$_2$              &  3.0 &  2.7 &   2.7 &  1.3  & 10.6  &   --  &  --   \\\\ \\\\ \\hline  \\end{tabular} \\end{table}  Our radiance ratio $7.2<E(4471)/E(5015)<9.8$ (Table\\,1) is in good agreement with $E(4471)/E(5015)=9.5$ from models with $T=8000$\\,K and $n_{\\rm H}\\approx 10^{10}$ cm$^{-3}$ by \\citet{heasley74}. However, the radiance ratio $41<E(4471)/E(4686)<62$ found in our study as well as in the eclipse data \\citep{sotirovski65, poletto67} is smaller than $E(4471)/E(4686)=6950$ obtained from those models. The latter ratio would predict for our faintest prominence (August\\,13 with $E(4471)=745$) a He\\,{\\sc ii}\\,4685.7\\AA{} radiance of $E(4686)\\approx0.1$ [erg~s$^{-1}$~cm$^{-2}$~sterad$^{-1}$] far below the noise level - even for eclipse observations. Those models are isothermal; recent models, however, with prominence--corona transition region, PCTR, show an increased He ionization (\\cite{labrosse04}). Our observed $E(4471)/E(4686)$ ratio is then in favor of an origin of He\\,{\\sc ii}\\,4685.7\\AA{} in the PCTR.   \\subsection{Line Widths}  A further hint of the PCTR contribution to the He lines is given by the different line widths observed for the three atomic states: singly ionized, triplet, singlet. The reduced width [$RW=\\Delta\\lambda_{\\rm D}/\\lambda$[ of the He\\,{\\sc ii}\\,4685.7\\AA{} line (after deconvolution of the atomic fine-structure) is significantly larger than that of the He\\,{\\sc i} triplet line: $1.18<RW(4686)/RW(4471)<1.8$, and the triplet is broader than the singlet line: $1.07<RW(4471)/RW(5015)<1.09$ (see Table\\,2). The latter result disagrees with observations by \\citet{heasley75}, bearing in mind their lower spectral resolution. Since the He\\,{\\sc i} lines are even visible in our raw spectra (see Fig.\\,1; in contrast to the faint He\\,{\\sc ii}\\,4685.7\\AA), the observed excess width $RW(4471)>RW(5015)$ is outside the error bars of at most $2\\%$ (see Table\\,2).  \\begin{table} \\caption{Reduced Doppler widths $\\Delta\\lambda_{\\rm D}/\\lambda$~[10$^{-5}$]; the values for He\\,{\\sc ii}\\,4685.7\\AA{} after deconvolution of the atomic fine-structure; numbers in parentheses give uncertainties originating from the fitting procedure.} \\begin{tabular}{lccccc}                       % \\hline emission line                 &   2\\,Aug.    &   4\\,Aug.    &  9\\,Aug.\\,a  & 9\\,Aug.\\,b  &   13\\,Aug.   \\\\ \\hline \\\\ He\\,{\\sc ii}\\,4685.7         & 5.64\\,(0.12) & 3.72\\,(0.13) & 4.21\\,(0.22) &      --     & 4.39\\,(0.39) \\\\ He\\,{\\sc i}\\,4471.5\\,(tripl) & 4.07\\,(0.01) & 3.14\\,(0.02) & 2.45\\,(0.01) & 2.21\\,(0.01)& 2.43\\,(0.01) \\\\ He\\,{\\sc i}\\,5015.7\\,(singl) &       --     &      --      & 2.24\\,(0.03) & 2.07\\,(0.04)& 2.28\\,(0.02) \\\\ Fe\\,{\\sc ii}\\,5018.4         &       --     &      --      & 1.53\\,(0.02) & 1.05\\,(0.01)& 1.32\\,(0.02) \\\\ Ti\\,{\\sc ii}\\,4468.4         & 3.28\\,(0.04) & 2.30\\,(0.04) & 1.17\\,(0.05) & 0.92\\,(0.02)& 1.21\\,(0.08) \\\\ Na\\,D$_2$\\,5890              & 3.17\\,(0.60) & 2.13\\,(0.03) & 1.36\\,(0.02) & 1.08\\,(0.01)& 1.20\\,(0.01) \\\\ \\\\  He\\,{\\sc ii}/tripl            &     1.37     &     1.18     &     1.72     &     --      &     1.80    \\\\ tripl/singl                   &       --     &       --     &     1.09     &     1.07    &     1.07    \\\\ tripl/Na\\,D$_2$           &     1.30     &     1.47     &     1.79     &     2.05    &     2.13    \\\\ \\\\  \\hline \\end{tabular} \\end{table}  The different widths of the He\\,{\\sc ii}, He\\,{\\sc i} triplet and singlet lines may be due to their formation in prominence regions of different temperature. Indeed, \\citet{stellmacher03} observe ''hotter'' lines to be more pronounced in such prominence regions which show less radiance in ''cooler'' lines. Labrosse (private communication, 2011) finds among 100 models with a PCTR of $10^5$K ({\\it cf.}, \\opencite{labrosse04}) mean ratios of the reduced widths of $1.1<RW(4686)/RW(4471)<1.6$ and $1.03<RW(4471)/RW(5015)<1.2$, respectively, in good agreement with our results. The PCTR then seems to be essential for an explanation of the the observed ''hierarchy'' of He line widths.  The difference between the singlet and the triplet emission is explained by the fact that the lowest He\\,{\\sc i} energy level belongs to the singlet system which can thus be populated directly by EUV radiation, while the triplet levels are populated from the (singlet) ground state mainly by ionization and recombination. Collisional excitation plays a minor role in cool prominence regions, which are well visible as dark absorption features in the Fe\\,XVI\\,335\\AA~image (see Figure\\,2). These cool (and dense) prominence cores emit the narrow Na\\,D lines which correspond to low T$_{\\rm kin}$ and $v_{\\rm nth}$ values as, e.g., the 7000\\,K and 2.8\\,km/s for August\\,13 in Figure\\,5.   \\begin{table} \\caption{Reduced Doppler width $\\Delta\\lambda_D/\\lambda$~[10$^{-5}$] and line radiance in absolute units [erg~s$^{-1}$~cm$^{-2}$~sterad$^{-1}$] of the metallic lines observed in the prominence at the east limb, $23^\\circ$N on Aug.13, 2011; radiance data are compared with \\citet{sotirovski65}.} \\begin{tabular}{cccccccccccc}                       % \\hline ion            &Sr\\,{\\sc ii}&Ca\\,{\\sc i}& Ti\\,{\\sc ii}& Fe\\,{\\sc ii}&Ba\\,{\\sc ii}&Ti\\,{\\sc ii}&Mg\\,{\\sc i}I&Na\\,{\\sc i}& Na\\,{\\sc i}\\\\ $\\lambda_0$\\,[\\AA] & 4215 & 4227 &  4468  &  4549  &  4554  &  4563  &  4571  & 5890  & 5896  \\\\ \\hline \\\\ reduced\\,width & 2.37 & 2.31 &  1.21  &  1.59  &  1.15  &  1.59  &  (0.87)&  1.20 & 1.14  \\\\ \\\\  obs. radiance   & 116  & 36.1 &  31.1  &  19.1  &  21.5  &  21.8  &   6.7  &   70  &  46   \\\\ Sotirovski     & 34.3 & 21.8 &  7.5   &   7.3  &   4.2  &   4.5  &   6.3  &  4.8  &  4.4  \\\\ \\\\ \\hline \\end{tabular} \\end{table}  The broad He lines, however, originate in hotter prominence regions where collisional excitation (and ionization) becomes effective. The observed ''hierarchy'' of excess broadening from He\\,{\\sc i} singlet to He\\,{\\sc i} triplet and He\\,{\\sc ii} lines may either be explained by differently hot threads or by a PCTR with outwards increasing temperature (or $v_{\\rm nth}$). But the width excess of He\\,{\\sc ii}\\,4685.7\\AA{} corresponds to 40000\\,K ({\\it cf.,} end of Section\\,3) which agrees with the formation temperature for He\\,{\\sc ii} lines calculated by \\citet{gouttebroze09}. Hence, a temperature increase through the PCTR, as assumed in the models by \\citet{heinzel01}, may be more realistic than an increase of $v_{\\rm nth}$. Also a combined outwards increase of both, T$_{\\rm kin}$ and $v_{\\rm nth}$, might fit our observations.  The singly ionized metallic lines show a similar width excess as He\\,{\\sc i}\\,4471\\AA{}: In Figure\\,5 their $v_0$ exceed $\\approx$1.2 times the values expected for [$T_{\\rm kin}=7000$\\,K; $v_{\\rm nth}=2.8$\\,km/s]. Hence, the singly ionized metallic lines may be emitted from similar layers as the He\\,{\\sc i} triplet line. The rather tight relation of radiance and widths of Na\\,D, He\\,{\\sc i}\\,4471\\AA{} and He\\,{\\sc ii}\\,4686\\AA{} (see Tables\\,1 and 2) favors a PCTR surrounding each individual thread rather than the prominence as a whole.  \\begin{acks} N.\\,Labrosse kindly put the model means at our disposal. We thank  an unknown referee for fruitful comments, J.\\,Hirzberger (MPS) for his code of fine-structure superposition and U.\\,Nolte (GWDG) for the figure layout. The research at IRSOL is financially supported by the Canton Ticino, the foundation Aldo e Cele Dacc\\`o, the city of Locarno, the local municipalities, and the SNF grant 200020-127329. R.\\,R. acknowledges the foundation Carlo e Albina Cavargna for financial support. \\end{acks}  \\newpage"
    },
    "1208/1208.2479_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "We have presented the key algorithmic developments that have gone into \\selavy, the prototype ASKAP source-finder. These features, including distributed processing, variable threshold determination and two-dimensional Gaussian profile fitting, have been implemented in a prototype system that has been made available to the ASKAP Survey Science Teams for testing purposes.  The development of the \\selavy source-finder is continuing as we move closer to ASKAP operations. Several Science Teams have provided feedback and specifications for additional or refined algorithms, covering pre-processing, source extraction and parameterisation and addressing some of the issues identified in this paper. At time of writing, we are incorporating these algorithms into \\selavy, with the aim of providing a more fully-fledged source-finder in time for BETA observations."
    },
    "1208/1208.2962_arXiv.txt": {
        "abstract": "In the cores of some galaxy clusters the hot intracluster plasma is dense enough that it should cool radiatively in the cluster's lifetime\\cite{lea73,cowie77, fabian77}, leading to continuous ``cooling flows'' of gas sinking towards the cluster center, yet no such cooling flow has been observed. The low observed star formation rates\\cite{odea08,mcdonald11b} and cool gas masses\\cite{edge01} for these ``cool core'' clusters suggest that much of the cooling must be offset by astrophysical feedback to prevent the formation of a runaway cooling flow\\cite{mcnamara07,fabian12,mathews09,gomez02}. Here we report X-ray, optical, and infrared observations of the galaxy cluster SPT-CLJ2344-4243\\cite{williamson11} at $z$ = 0.596. These observations reveal an exceptionally luminous ($L_{2-10~{\\rm keV}} = 8.2\\times10^{45}$ erg s$^{-1}$) galaxy cluster which hosts an extremely strong cooling flow (\\.{M}$_{cool}$ = 3820 $\\pm$ 530 M$_{\\odot}$ yr$^{-1}$). Further, the central galaxy in this cluster appears to be experiencing a massive starburst (740 $\\pm$ 160 M$_{\\odot}$ yr$^{-1}$), which suggests that the feedback source responsible for preventing runaway cooling in nearby cool core clusters may not yet be fully established in SPT-CLJ2344-4243. This large star formation rate implies that a significant fraction of the stars in the central galaxy of this cluster may form via accretion of the intracluster medium, rather than the current picture of central galaxies assembling entirely via mergers. ",
        "introduction": "\\setcounter{figure}{0} \\renewcommand*\\thefigure{S.\\arabic{figure}} \\setcounter{table}{0} \\renewcommand*\\thetable{S.\\arabic{table}} \\subsection{X-ray Data:} Our data reduction pipeline was adapted from earlier work\\cite{vikhlinin05, A11} and includes the removal of flares, estimation of background from blank sky fields, and calibration using the latest set of corrections. From the cleaned image, an X-ray spectrum was extracted over the energy range 0.5--8.0 keV within an annulus defined as 0.15r$_{500}$ $<$ r $<$ r$_{500}$, excluding point sources, with the initial assumption that r$_{500}$ = 1.0 Mpc. We fit this spectrum using a combined WABS(MEKAL) model, which accounts for Galactic absorption in the emission spectrum of a hot, diffuse gas. This fit yields a temperature (T$_X$) and gas mass (M$_g$), which are combined to give the mass proxy Y$_X$ $\\equiv$ M$_g$ $\\times$ T$_X$.  Assuming the following scaling relation\\cite{vikhlinin09a}: \\begin{equation} M_{500} = 5.77\\times10^{14} h^{1/2} M_{\\odot} \\times \\left(\\frac{Y_X}{3\\times10^{14} M_{\\odot} keV}\\right)^{0.57} E(z)^{-2/5}~, \\end{equation} we can then infer M$_{500}$, which leads to a new value of r$_{500}$ based on its definition: \\begin{equation} r_{500} \\equiv \\left(\\frac{3M_{500}}{4\\pi500\\rho_{crit}(z)}\\right)^{1/3} . \\end{equation} This process is iterated until changes in $r_{500}$ are small. From the final estimate of M$_{500}$, we can extrapolate M$_{200}$ (which, by convention, uses the average density of the Universe, rather than the critical density) assuming that the dark matter halo has an NFW profile with a concentration following the mass-concentration relation\\cite{duffy08}. At this point, we measure T$_X$ within the aperture 0.15r$_{500}$ $<$ r $<$ r$_{500}$, and L$_X$ and M$_g$ within the aperture 0 $<$ r $<$ r$_{500}$, which are used as the final global quantities.  To determine an inner characteristic radius, we compute the hard (2.0--8.0 keV) and soft (0.7--2.0 keV) X-ray surface brightness profiles, as shown in Figure \\ref{sbprof}. From this plot, it is clear that the central $\\sim$10 kpc is dominated by a hard X-ray point source, likely an AGN, and then the hardness ratio reaches a minimum over the range 10--100 kpc, before settling to a roughly constant value from 100--1300 kpc. This inner region with excess soft X-ray emission defines the cool core, with a cooling radius of $\\sim$100 kpc. Further, the right panel of Figure \\ref{sbprof} shows that the classical cooling rate does not increase substantially by including emission from r $>$ 100 kpc, suggesting that the cool core is confined to r $<$ 100kpc.  \\begin{figure*}[t] \\begin{center} \\begin{tabular}{cc} \\includegraphics[width=0.45\\textwidth]{sbprof.pdf} & \\includegraphics[width=0.48\\textwidth]{dmdt_r.pdf} \\\\ \\end{tabular} \\caption{Left: X-ray emission measure as a function of radius for SPT-CLJ2344-4243. The red and blue lines show the contributions to the total (black) from soft (0.7--2.0 keV) and hard (2.0--8.0 keV) emission in the observed frame, respectively. The lower panel shows the hardness ratio, HR = $\\frac{H-S}{H+S}$ as a function of radius. The spectrally hard, spatially unresolved AGN emission dominates at r $<$ 10 kpc, while the emission from 10 kpc $<$ r $<$ 100 kpc has a soft X-ray excess, which we interpret as a cool core. Note that the error bars are smaller than the point size in this radial range. Right: Classical cooling rate as a function of enclosed radius. This plot shows that the cooling rate rises rapidly out to $\\sim$100kpc, at which point it changes little out to $>$200kpc, due to the fact that gas at large radii has a much longer cooling time. This plot further motivates our choice of a 100kpc cooling radius.} \\label{sbprof} \\end{center} \\end{figure*}   We extracted spectra in the logarithmically-spaced annuli 0 $<$ r $<$ 100 kpc (cool core), 100 kpc $<$ r $<$ 450 kpc, and 450 kpc $<$ r $<$ 1300 kpc  in order to determine the gas temperature (T$_X$), electron density (n$_e$), and metallicity (Z) as a function of radius. In the central bin, we masked the inner 1.5$^{\\prime\\prime}$ in order to remove contributions to the spectrum from the AGN. These spectra were fit with a model combining Galactic absorption (WABS) and hot, diffuse gas (MEKAL).  The specific entropy (K = T$_X$ $\\times$ n$_e^{-2/3}$) and cooling time (t$_{cool}$ = 10$^8$ $\\left(\\frac{K^{3/2}}{10}\\right)$$\\left(\\frac{T_X}{5}\\right)^{-1}$ Gyr) were also inferred from the model fit to the spectrum in each annulus.   \\begin{figure*}[p] \\begin{center} \\includegraphics[width=0.6\\textwidth]{xray_profs.pdf} \\caption{Temperature, metallicity, specific entropy, and cooling time profiles for SPT-CLJ2344-4243. The outer radius (1300 kpc) corresponds to $r_{500}$ for this system, while the inner radius (100 kpc) corresponds to the size of the cool core. The dip in temperature and peak in metallicity in the central bin is reminiscent of $z\\sim0$ cool core clusters. In the central 100 kpc, we find that, without correcting for projection, the cooling time is $<$ 1 Gyr, suggesting a strong cooling flow.} \\label{xrayprofs} \\end{center} \\end{figure*}  The projected temperature and metallicity profiles show a rise in metallicity accompanied by a dip in temperature in the central 100 kpc, as seen in cool core clusters at z $\\sim$ 0. Despite the fact that these data have not been corrected for projection, which means that a significant amount of emission from hot gas at large radius is contributing to the spectrum extracted from the central aperture, the cooling time in the inner 100 kpc is $<$ 1 Gyr. This short central cooling time, along with the low ($<$100 keV cm$^2$) central entropy, resembles nearby strong cool cores, such as the Perseus and PKS0745-191 clusters.  The classical cooling rate was estimated from the equation, $\\frac{dM}{dt}$ = $\\frac{2L\\mu m_p}{5kT}$, using the total X-ray luminosity and temperature within the central 100 kpc\\cite{odea08}. As shown in Figure \\ref{sbprof}, this estimate is relatively insensitive to the choice of radius for r $\\geq$ 100 kpc. We find dM/dt$_{classical}$ = 3820 $\\pm$ 530 M$_{\\odot}$ yr$^{-1}$. The available X-ray spectra are of insufficient depth to estimate the cooling rate spectroscopically and thus we are unable to measure the ``instantaneous'' cooling rate. However, the classical, steady-state, cooling rate is still valuable for comparison to low-$z$ clusters and as an upper limit to the ``true'' cooling rate.  \\subsection{Infrared Imaging} Snapshot Herschel PACS images were obtained at 70 and 160$\\mu$m with a total exposure time of $<$30 minutes as part of a Director's Discretionary Time project (PI: M. Bayliss). In these short exposures, the central galaxy was detected at flux levels of 432 $\\pm$ 21 mJy and 364 $\\pm$ 18 mJy, at 70 and 160$\\mu$m, respectively.  Herschel SPIRE maps at 250, 350 and 500$\\mu$m were observed as part of the ``Herschel Lensing Survey'' (HLS; PI: E. Egami) snapshot program covering 148 SPT clusters. The SPIRE data consists of a single repetition map, with coverage complete to a cluster-centric radius of ~5 arcmin. The maps were produced via the standard reduction pipeline HIPE v9.0, the SPIRE Photometer Interactive Analysis (SPIA) package v1.7, and the calibration product v8.1, with improved treatment of the baseline removal (also known as `de-striping').  To characterize the far-IR SED, we fit a blackbody law, modified with a spectral emissivity that varies physically such that the dust opacity reaches unity at frequency $\\nu_{\\rm c}$ \\cite{blain03}: \\begin{equation} f_{\\nu}\\propto [1-{\\rm exp}(-(\\nu / \\nu_{\\rm c})^\\beta)]B_{\\nu}(T_{\\rm d}) \\end{equation} Here, $B_{\\nu}(T_{\\rm d})$ is the Planck function. We fix the spectral index of the emissivity to $\\beta = 2.0$, and the critical frequency to $\\nu_{\\rm c} = 1.5$ THz.  The dust temperature $T_d$ and the amplitude are left as free parameters. We exclude photometric data at wavelengths shorter than rest wavelength $\\sim 40\\,{\\rm \\mu m}$ (including 70, 160, 250, 350, 500 $\\mu$m data from Herschel PACS and SPIRE) so we can fit only to the cold dust component, which should be more free of AGN contamination and better trace the dust heated by star formation. Our best fit gives $T_d=87 \\pm 3 $\\ K and $L_{IR}=(9.5 \\pm 1.1)\\times 10^{12} ~L_{\\odot}$, with $r\\chi ^2=0.13$.  \\subsection{Optical Spectroscopy:} Long-slit optical spectra at two different orientations (104$^{\\circ}$ and 135$^{\\circ}$) with a 1.2$^{\\prime\\prime}$ slit width were obtained using the IMACS spectrograph on the Baade 6.5m telescope, using the 200 lines/mm grism, which provides 2.0\\AA/pixel spectral resolution over the wavelength range 3900\\AA--10000\\AA. The seeing during these observations was $\\sim$0.7$^{\\prime\\prime}$. These spectra were reduced using standard IRAF tasks to remove the bias and overscan, flat field, remove sky lines, and perform wavelength calibration based on arc lamp spectra. The LA Cosmic software\\cite{lacosmic} was used to mask cosmic rays before combining exposures. Flux calibration was performed by measuring the $g$, $r$, $i$, $z$ flux within a 1.2$^{\\prime\\prime}$$\\times$1.2$^{\\prime\\prime}$ box centered on the central galaxy nucleus from the broadband imaging, and forcing the spectrum, extracted from a boxcar with the same spatial and spectral dimensions, to pass through these points. The final reduced spectrum in a 1.2$^{\\prime\\prime}$$\\times$1.2$^{\\prime\\prime}$ extraction region is shown in Figure \\ref{optsed}.   \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.6\\textwidth]{optsed.pdf} \\caption{IMACS spectrum of the central galaxy in SPT-CLJ2344-4243. This spectrum demonstrates the strength of the nebular emission lines ([O II], H$\\beta$, [O III]), as well as the relative flatness of the continuum spectrum around the 4000\\AA\\ break, indicating a relatively young population of stars. The red circles represent broadband fluxes, extracted in the same region is the spectrum, which were used to flux-calibrate the spectra. } \\label{optsed} \\end{center} \\end{figure*}  The velocity dispersion for SPT-CLJ2344-4243 is estimated from cluster member galaxy recession velocities that were measured using the Gemini Multi-Object Spectrograph (GMOS) as a part of a large NOAO survey program (PI: C. Stubbs) to measure velocity dispersions for 100 SPT galaxy clusters. Galaxies were prioritized for MOS slits based on proximity to the red-sequence and magnitude. The raw spectra were bias-subtracted, flat-fielded, wavelength calibrated, and mapped to a common mosaic grid using the \\emph{gemini.gmos} IRAF package. The reduced 2D spectral exposures were sky-subtracted, extracted, stacked, and flux calibrated relative to LTT 1788 using custom IDL routines that make use of the XIDL package\\footnote{http://www.ucolick.org/\\textasciitilde xavier/IDL/}. Velocity measurements were made using the RVSAO\\cite{kurtz98} package with the \\emph{fabtemp97} template. The final histogram of velocities, along with the biweight estimate of the redshift and velocity dispersion, are shown in Figure \\ref{velhist}.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.6\\textwidth]{velhist.pdf} \\caption{Radial velocities of 26 galaxies in SPT-CLJ2344-4243, relative to the mean cluster redshift of $<z>$ = 0.596. The best-fit Gaussian is shown as a solid line. There is some evidence for substructure at negative velocities, but the limited number of redshifts do not allow for a statistically significant result.} \\label{velhist} \\end{center} \\end{figure*}  \\subsection{Near-Infrared Spectroscopy} A near-infrared spectrum of SPT-CLJ2344-4243 was obtained with the Folded-port Infrared Echellette (FIRE) spectrograph at the Magellan Baade telescope in late January, 2012. FIRE delivers $R=6000$ spectra between $0.82-2.5$ microns in a single-object, cross-dispersed setup\\cite{sim10}. A 10-minute on-target exposure with $\\sim$0.9$^{\\prime\\prime}$ seeing produced the data presented in Figure \\ref{firespec}.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=0.6\\textwidth]{fire.pdf} \\caption{IMACS optical (blue) and FIRE infrared (red) spectra of the central galaxy in SPT-CLJ2344-4243. The IMACS spectrum has been scaled down to account for the significantly shorter and narrower FIRE slit.} \\label{firespec} \\end{center} \\end{figure*}  For the extraction of point sources, FIRE's reduction pipeline (FIREHOSE) nominally creates a 2-dimensional sky model derived from the portions of the slit that are not illuminated by the source. This way, the sky flux is measured simultaneously with the object flux. Since the spatial extent of SPT-CLJ2344-4243 fills FIRE's $6^{\\prime\\prime}$-long echelle slit, a separate 10-minute sky exposure had to be obtained and subtracted from the object frame prior to extraction of the spectrum. The variability in the sky over the fifteen minutes between the object and sky exposures can introduce uncertainties into the extracted object spectrum, particularly near hydroxyl (OH) lines. Fortunately, the OH lines subtract with few residuals in the vicinity of the H$\\alpha$ emission.  Flux calibration was performed by obtaining the spectrum of an A0V star with an airmass, angular position, and observing time as close to the target as possible. Telluric absorption was corrected\\cite{vacca03} via the xtellcor procedure released with the spextool pipeline\\cite{cushing04}.  Since FIRE operates in quasi-Littrow mode, the spectral orders are significantly curved and tilted with respect to the detector's pixel basis. Hence, producing a spatial by spectral image of SPT-CLJ2344-4243 required a separate boxcar extraction of the object spectrum for each spatial position. A boxcar width of $0.2^{\\prime\\prime}$ was selected to match the spatial scale of IMACS. Each of the spectral strips produced from these extractions were telluric-corrected with the observation of an AOV standard. ",
        "conclusions": ""
    },
    "1208/1208.6400_arXiv.txt": {
        "abstract": "Non-equilibrium radiation diffusion is an important mechanism of energy transport in Inertial Confinement Fusion, astrophysical plasmas, furnaces and heat exchangers. In this paper, an analytical solution to the non-equilibrium Marshak diffusion problem in a planar slab and spherical shell of finite thickness is presented. Using Laplace transform method, the radiation and material energy densities are obtained as a function of space and time. The variation in integrated energy densities and leakage currents are also studied. In order to linearize the radiation transport and material energy equation, the heat capacity is assumed to be proportional to the cube of the material temperature. The steady state energy densities show linear variation along the depth of the planar slab, whereas non-linear dependence is observed for the spherical shell. The analytical energy densities show good agreement with those obtained from finite difference method using small mesh width and time step. The benchmark results obtained in this work can be used to validate and verify non equilibrium radiation diffusion computer codes in both planar and spherical geometry. ",
        "introduction": "The time dependent non-equilibrium radiation transport equation is non linearly coupled to the material energy equation\\citep{Pomraning-book},\\citep{Mihalas-book}. Also the material properties have complex dependence on the independent variables. As a result, the time dependent thermal radiation transport problems are commonly solved numerically. Several numerical methods are in use for this purpose, namely the discrete ordinates \\citep{Ghosh}, finite volume \\citep{Kim}, Monte Carlo \\citep{Fleck}, hybrid stochastic-deterministic \\citep{Densmore},\\citep{Connolly}, or the approximate methods like the Eddington approximation \\citep{Shettle}, heat conduction \\citep{Goldstein} or the diffusion approximations \\citep{Dai}, \\citep{Knoll}, \\citep{Ober}. Benchmark results for test problems are necessary to validate and verify the numerical codes \\citep{Ensman}. Analytical solutions producing explicit expressions for the radiation and material energy density, integrated densities, leakage currents, etc. are the most desirable.  In the literature, considerable amount of efforts have been applied for solving the Radiation Transport problem analytically. Marshak obtained a semi-analytical solution by considering radiation diffusion in a semi infinite planar slab with radiation incident upon the surface \\citep{Marshak-original}. Assuming that the radiation and material fields are in equilibrium, the problem admits a similarity solution to a second order ordinary differential equation which was solved numerically \\citep{Kass}. The results were extended for non-equilibrium radiation diffusion by assuming that the specific heat is proportional to the cube of the temperature \\citep{Pomraning}, \\citep{Su-olson1}. This assumption linearized the problem providing a detailed analytical solution. As the radiative transfer codes are meant to handle an arbitrary temperature dependence of the material properties, the obtained  solutions serve as a useful test problem \\citep{Ganapol-pomraning}, \\citep{Su-olson2}, \\citep{Su-olson3}. Using the same linearization, 3T radiation diffusion equations were solved for spherical and spherical shell sources in an infinite medium \\citep{MCClarren}. All available results on the non-equilibrium radiative transfer problems in planar and spherical geometry consider systems having infinite or semi-infinite extension. Benchmarks involving finite size systems have been limited either to the heat conduction or equilibrium diffusion approximation \\citep{Williams},  \\citep{Olson-H}, \\citep{Liemert}.  In this paper, we solve the time dependent non-equilibrium radiation diffusion problem for finite size systems in both planar and spherical geometry. Non-equilibrium diffusion codes can be more easily validated and verified against these benchmark results because there is no need to consider a slab or spherical medium of  very large size for avoiding boundary effects. Analytical expressions for all the quantities of interest can be obtained for finite slab/shell width and parameter values relevant to practical problems. This work can be extended to multi-dimension using separation of variables and Laplace transform method or the eigenfunction expansion method to obtain analytical series solution in a manner similar to the multilayer heat conduction\\citep{Jain}.  The remainder of the paper is organised as follows. In Section \\ref{analytic}, the analytical solution for the finite planar slab and spherical shell is derived followed by Section \\ref{numerical} on numerical finite difference method. In Section \\ref{results}, the results for the radiation and material energy densities, leakage currents, integrated quantities, etc. are plotted and physically explained. Finally, conclusions are given in Section \\ref{conc}. \\label{intro} ",
        "conclusions": "\\label{conc} In this paper, the time dependent non equilibrium radiation diffusion problem has been solved analytically for finite planar slab and spherical shell with a constant radiation flux incident on the surface. The observed trend in temporal and spatial variation of energy densities, leakage currents, integral quantities, etc. has been explained physically. The analytical values of energy densities are cross checked with the solution of finite difference analysis and good agreement is observed when small mesh width and time steps are used. The results obtained in this paper can serve as new and useful benchmarks for non equilibrium radiation diffusion codes in both planar and spherical geometries. The same methodology can be applied to any other finite size systems like layered media with various boundary conditions. Using separation of variables, the method can be extended to generate benchmark results for validating radiation diffusion codes in two and three dimensions. \\\\  {\\bf Acknowledgement:} \\\\ The author would like to thank Dr. N. K. Gupta for his useful comments, support and encouragement."
    },
    "1208/1208.3539_arXiv.txt": {
        "abstract": "Reconstruction of the curvatures of radio wavefronts of air showers initiated by ultra high energy cosmic rays is discussed based on minimization algorithms commonly used. We emphasize the importance of the convergence process induced by the minimization of a non-linear least squares function that affects the results in terms of degeneration of the solutions and bias. We derive a simple method to obtain a satisfactory estimate of the location of the main point of emission source, which mitigates the problems previously encountered. ",
        "introduction": "The determination of the nature of the ultra-high energy cosmic rays (UHECR) is an old fundamental problem in cosmic rays studies. Numerous are the difficulties. New promising approaches could emerge from the use of the radio-detection method which exploits, through antennas, the radio signal that accompanies the development of the extensive air shower (EAS). Several experimental prototypes like CODALEMA \\cite{key-1} in France and LOPES \\cite{key-2} in Germany shown the feasibility and the potential of the method to reconstruct EAS parameters, as the arrival direction, the impact location at ground, the lateral distribution function of the electric field, or the primary particle energy \\cite{key-3,key-4,key-5,key-6,key-7}. However, the temporal radio wavefront characteristics remain still poorly determined \\cite{key-8,key-8-1}, although its knowledge could be consider as one of the first steps in retrieving information about the EAS itself. The importance of this information resides in its potential sensitivity to the nature of the primary particle, especially because the existence of a curvate radio wavefront (a spherical wavefront) could provide the location of the main point of the emission source, and possibly an estimation of Xmax, event by event. Indeed, the arrival timing being defined by the maximum amplitude of the radio signal, it is more likely linked to a limited portion of the longitudinal development of the shower (and so especially at the point of maximum) \\cite{key-8-2}.\\medskip{}   Moreover, the migration of present small scale radio-prototypes to large scale experiments spread over surfaces of several tens of $1000\\: km^{2}$ using self-triggered antennas, is challenging. This technique is subjected to delicate limitations in regard to UHECR recognition, due to noises induced by human activities (high voltage power lines, electric transformers, cars, trains and planes) or by stormy weather conditions (lightning). Figure \\ref{codanoise} shows a typical reconstruction of sources obtained with the CODALEMA experiment \\cite{key-9}, by invoking a spherical wave minimization. Such patterns are also commonly observed in others radio experiments \\cite{key-9-1,key-9-2}. In most of the cases, one of the striking results is that these emission sources are reconstructed with great inaccuracy, although they are fixed and although the number of measured events is high. By extension, a cosmic event being a single realization of the detected observables (arrival time and peak amplitude on each antenna), interpretation of such methods of reconstruction for the identification of a point source can become even more delicate, even using statistical approaches.  \\begin{figure}[hbtp] \\centering \\includegraphics[width=11cm,height=6cm]{BKG_2011_06_21_source2} \\caption{Typical result of reconstruction of two entropic emitters at ground, observed with the stand-alone stations of CODALEMA, through standard minimization algorithms. Despite the spreading of the reconstructed positions, these two transmitters are, in reality, two stationary point sources. } \\label{codanoise} \\end{figure}   The commonly used technique relies on the minimization of an objective function which depends on the assumed shape of the wavefront, using the arrival times and locations of the antennas. The aim of this paper is to highlight that the minimization of such an objective function, incorporating a spherical wave front, can be an ill-posed problem. We will show that it originates from strong dependencies of the convergence of the minimization algorithms with initial parameters, from the existence of degenerations of the solutions (half lines) which can trap most of the common algorithms, and from the existence of offsets (bias) in the reconstructed positions. Finally, by avoiding more complex estimates based on advanced statistical theories, we got to deduce a simple method to obtain a significant estimate of the source location. We compared the exact results with our numerical reconstructions performed on a test array. ",
        "conclusions": "Experimental results indicated that the common methods of minimization of spherical wavefronts could induce a mis-localisation of the emission sources. In the current form of our objective function, a first elementary mathematical study indicates that the source localization method may lead to ill-posed problems, according to the actual source position. To overcome this difficulty, we developed a simple method, based on grid calculation of the objective function. This approach appears to provide, at worst, an estimate as good as for the common algorithms for locating the main point of the emission source, keeping in mind that this method is not optimal in the sense of optimization theories. However, further developments are without any doubt still necessary, maybe based on advanced statistical theories, like for instance by adding further information (as the signal amplitude or the functional of the radio lateral distribution). This could be achieved by trying a generalized objective function which includes these parameters. In addition, the interactions with other disciplines which face this problem could also provide tracks of work (especially regarding earth sciences which focus on technics of petroleum prospecting)."
    },
    "1208/1208.1827_arXiv.txt": {
        "abstract": "Vela X-1 is the archetype of high-mass X-ray binaries, composed of a neutron star and a massive B supergiant. The supergiant is a source of a strong radiatively-driven stellar wind. The neutron star sweeps up this wind, and creates a huge amount of X-rays as a result of energy release during the process of wind accretion. Here we provide detailed NLTE models of the Vela X-1 envelope. We study how the X-rays photoionize the wind and destroy the ions responsible for the wind acceleration. The resulting decrease of the radiative force explains the observed reduction of the wind terminal velocity in a direction to the neutron star. The X-rays create a distinct photoionized region around the neutron star filled with a stagnating flow. The existence of such photoionized bubbles is a general property of high-mass X-ray binaries. We unveiled a new principle governing these complex objects, according to which there is an upper limit to the X-ray luminosity the compact star can have without suspending the wind due to inefficient line driving. ",
        "introduction": "A high-mass X-ray binary (HMXB) is a binary star system consisting of a massive luminous hot star (frequently OB supergiant) and a compact object, either a neutron star or a black hole \\citep{rem}. A fraction of the stellar wind of the luminous hot star is trapped in the gravitational well of the compact object, and is accreted onto its surface \\citep{davos,laheupet}. Part of the released potential energy of accreting material is transformed into X-rays, resulting in one of the most powerful stellar X-ray sources. Such systems of stars in interaction belong to the most valuable astrophysical laboratories. The binary nature of the object enables us to determine stellar parameters precisely, which subsequently serve as a firm base for further study.  Vela X-1 (HD\\,77581, GP\\,Vel) is the archetype of high-mass X-ray binaries, consisting of a neutron star and a massive B supergiant \\citep{chodil,brucato,barziv,tom}. The neutron star is a source of pulsed X-ray and $\\gamma$-ray emission with a period of 283\\,s \\citep{mekac,sever}, modulated both by the orbital motion and stochastic variations \\citep{bildik}. The X-rays propagating through the hot star wind probe the wind structure, yielding information about the mass-loss rate and the velocity field \\citep{viteal}. The perpetual X-ray variation (flaring) reveals the existence of some structure in the wind -- clumping \\citep{ducim,prvni}. On the other hand, X-rays also significantly influence the stellar wind, resulting in X-ray photoionization of its material \\citep{mav,viteal}. Because the stellar wind of hot stars is mostly driven by the light absorption in the lines of heavier elements, the X-ray photoionization may influence the wind acceleration. Particularly, the appearance of highly charged ions, which absorb the light less effectively than low-charged ions, causes the decrease of the radiative force. Since this force is responsible for driving the wind, the wind flow may subsequently stagnate.  Numerical studies of stellar winds in HMXBs concentrate mainly on the multidimensional simulation of wind accretion \\citep{blok,blowoo,felan,hadvitr}, while the wind driving is simplified using force multipliers that take the X-ray irradiation into an account in an approximative way \\citep{stekal,stesam}. This is a significant shortcoming, because the X-ray photoionization affects the radiative force, and consequently the amount and velocity of wind material accreted on the compact companion. Detailed modelling of ionization and excitation balance in the wind is crucial for the understanding of the influence of X-ray photoionization on the wind dynamics. The ionization and excitation balance should be properly derived using equations of statistical equilibrium.% \\footnote{% This approach is usually referred to as non-LTE or NLTE and it means that the assumption of thermodynamic equilibrium is not used for evaluation of the excitation and ionization balance.} To remedy the situation, we provide wind models that include the influence of X-ray irradiation using up-to-date NLTE models. ",
        "conclusions": "We provide detailed numerical models of the influence of X-rays on the supergiant wind in the Vela X-1 binary system. The effect of X-ray photoionization on the radiative force and wind dynamics has never been studied using appropriate NLTE wind models. The X-rays photoionize the wind and destroy the ions responsible for the wind acceleration. This results in flow stagnation in the vicinity of the neutron star, which was identified in observations. For a sufficiently strong X-ray source the wind that directly faces the neutron star falls back on the mass losing star and never reaches the compact companion.  We have shown that there is an upper limit to the X-ray luminosity the compact star can have without disrupting the stellar wind. For a higher luminosity than the limiting one the decrease of the wind acceleration is so strong that no wind material would reach the neutron star. This theoretical picture of the maximum X-ray luminosity is supported by observation of many high-mass X-ray binaries.  The wind equations allow the existence of two types of solutions with different X-ray luminosities and wind velocities. The case of a strong X-ray source, which significantly affects the wind ionization, leads to accretion of slow wind in large amounts resulting in a strong X-ray source. On the other hand, a weak X-ray source that does not significantly influence the wind ionization results in accretion of fast wind in low amounts producing a weak X-ray source. Different types of solution may appear in different binary systems, or perturbations may cause switching between these two types of wind solutions contributing to the X-ray variability.  The predicted mass-loss rate agrees with the value estimated from X-ray spectroscopy. Moreover, the wind mass-loss rate cannot be lower than this value, because a decrease of the wind mass-loss rate would lead to the disruption of the wind and disappearance of the X-ray emission. This supports the reliability of current mass-loss rate predictions."
    },
    "1208/1208.4775_arXiv.txt": {
        "abstract": "We study the vertical extent of propeller structures in Saturn's rings. Our focus lies on the gap region of the propeller and on non-inclined propeller moonlets. In order to describe the vertical structure of propellers we extend the model of \\citet{spahn00} to include the vertical direction. We find that the gravitational interaction of ring particles with the non-inclined moonlet does not induce considerable vertical excursions of ring particles, but causes a considerable thermal motion in the ring plane. We expect ring particle collisions to partly convert the lateral induced thermal motion into vertical excursions of ring particles. For the gap region of the propeller, we calculate gap averaged propeller heights on the order of 0.7 Hill radii, which is of the order of the moonlet radius. In our model the propeller height decreases exponentially until viscous heating and collisional cooling balance. We estimate Hill radii of $370$m and $615$m for the propellers Earhart and Bl\\'eriot. Our model predicts about $120$km for the azimuthal extent of the Earhart propeller at Saturn's 2009 equinox, being consistent with values determined from Cassini images. ",
        "introduction": "\\label{sec:introduction}  Planetary rings are one of the most remarkable and beautiful cosmic structures. They are considered to be natural dynamical laboratories \\citep{burns06}, exemplifying the physics of cosmic disks, such as accretion disks or galactic disks, which are much larger and much farther away from Earth. An exciting example is the presence of small moons embedded in Saturn's rings, henceforth called moonlets, which have their analog in planetary embryos orbiting within a protoplanetary disk \\citep{artymowicz06,papaloizou07}.  The fact that the resolution of even \\emph{Cassini's} cameras is too low to image these moonlets directly, brought up the idea of investigating moonlet induced putative structures in the rings \\citep{spahn87,spahn89}, with the hope that these features could be captured by the spacecraft cameras or instruments. This then led to predictions of the propeller structures \\citep{spahn00,sremcevic02} which are carved in the rings by the moonlet. Subsequent numerical particle experiments \\citep{seiss05,sremcevic07,lewis09} completed the \\emph{fingerprint} of such gravitational perturbers and confirmed the spatial scaling of the propeller structure. Depending on its size, an embedded ring-moon either creates a propeller (sizes below 1 km) or, alternatively, opens up a circumferential gap (for sizes $>$ 1 km, e.g. the ring-moons {\\em Pan} and {\\em Daphnis}). Both structures, propeller and circumferential gap, are decorated with density wakes, completing the structural picture.  Up to this stage, all these density features have been assumed to occur only in the ring plane, a vertical stratification of moonlet induced structures has not seemed to be of importance.  More than 150 propeller moonlets have now been detected \\citep{tiscareno06,tiscareno08} and among them a few which are large enough to allow \\emph{Cassini's} cameras to take several snapshots of their propellers at different times, confirming in this way their orbital motion. Those moonlets were nicknamed after famous aviators, e.g.: Bl\\'{e}riot, Kingsford Smith, Earhart \\citep{tiscareno10}.  In the summer of 2009, at Saturn's equinox (the Sunset at Saturn's rings), the perfect opportunity arose to detect any vertical structure deviating from the mean ring plane by observing shadows cast on the rings. At this time the density structures around the largest propeller moonlets, as well as those around the ring-moon \\emph{Daphnis}, created prominent shadows. These can be assigned to the wakes and in the case of the propeller moonlets also to excited regions of the propeller, where the moonlet induces a partial gap. The shadows were much longer than the moon's size itself. The lateral extent of the shadows allows to conclude that moonlet induced vertical excursions of ring particles can be in the range of several kilometers in the case of \\emph{Daphnis} or several hundred meters in the case of the large propeller moonlets. These very facts directly indicate the necessity to investigate the vertical stratification of moonlet induced structures, which has not been the focus of former models of the moonlet's \\emph{fingerprint}. In this work we will study the vertical extent of moonlet induced propeller stuctures, focusing on the gap region of the propeller.  The paper is organized as follows: In Section \\ref{sec:extended_model} the extended propeller model is presented. In Section \\ref{sec:gravitational_scattering} the mass flow through the scattering region is calculated by a probabilistic approach, and values of the moonlet induced thermal velocities are determined, which are later used as initial conditions for the hydrdynamical equations. Section \\ref{sec:hydrodynamic_flow} gives the hydrodynamical balance equations, which we use to model the diffusion of mass into the induced gap and the relaxation of the ring temperature. In Section \\ref{sec:propeller_height} we calculate the height of the ring in the gap region of a propeller. Section \\ref{sec:discussion} discusses made assumptions and the application of our results to propeller features in Saturn's rings. Our results are summarized in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:discussion}  We study the vertical extent of propeller structures in Saturn's rings, focusing on the propeller gap region. The effective geometric thickness $H_\\text{eff}$ is used to describe the propeller height as a function of the ring temperature, calculated using the equilibrium ratios of the thermal velocities. A vertically constant ring temperature serves as a fair assumption, because for the low optical depths in the propeller gap region the vertical dependence of the ring temperature is rather weak \\citep{schmidt99}.  In our model the azimuthal temperature decrease is caused dominantly by the disturbed balance of viscous heating and collisional cooling. We assume a constant coefficient of restitution for the cooling term and constant viscosity for the viscous heating term. This is a simplification, but allows a semi-analytical solution of the azimuthal temperature relaxation (some integrals have to be numerically evaluated), and most important, it is consistent with images of propeller shadows taken by the Cassini spacecraft.  The restriction of a constant $\\nu_0$ could be dropped by introducing e.g. a power law dependence \\citep{spahn00b}, \\begin{equation} \\nu = \\nu_0 \\left(\\frac{\\Sigma}{\\Sigma_0}\\right)^\\beta \\left(\\frac{r}{a_0}\\right)^\\gamma \\left(\\frac{T}{T_0}\\right)^\\alpha\\ . \\label{eq:viscosity} \\end{equation} For small moonlets ($h \\le 50$m) one has a rather small ratio $T_\\text{ini}/T_0 \\le 5$. The power law dependence (\\ref{eq:viscosity}) with $\\alpha = 1/2$ \\citep{salo01} then gives a viscosity about twice as large as the unperturbed one. On the other hand, the surface mass density in the gap region after the scattering by the moonlet is about half of the unperturbed value, so that a constant viscosity $\\nu_0$ can be justified.  In case of large moonlets the ratio $T_\\text{ini}/T_0$ is by far higher ($190$ for a moonlet with $300$m Hill radius at $|\\tilde{x}| = 2$). The viscosity will rise considerably with increasing temperature, tending to an increased heating, and thus, to an increased cooling time scale. However, because of an increasing coefficient of restitution and due to increased collision rates, an enhanced cooling works against that effect. Insofar, as a first step, we have used a constant viscosity and a constant coefficient of restitution, interpreting both as effective values.  Furthermore, we have neglected the temperature decrease due to heat conduction, based on the grounds that the relaxation of the ring temperature by granular cooling is the faster process, compared to the large diffusive time scales blurring the propeller gaps. However, heat conduction as well as viscous diffusion could become important for small moonlets. We have estimated the temperature decrease due to heat conduction for a simplified model, described in \\ref{app:heat_conduction}. For small moonlets ($h \\le 50$m) the temperature decrease due to heat conduction is comparable to the temperature decrease of our granular cooling driven solution (\\ref{eq:tsolution}). On the other hand, for large moonlets the temperature decrease due to heat conduction was approximatly $20$ times smaller than our cooling driven relaxation after $10$ orbits, justifying a neglect of heat conduction for large moonlets, and thus, our approach.   \\begin{figure} \\begin{adjustwidth}{-1in}{-1in}% \\begin{subfigure}[t]{0.7\\textwidth} \\centering \\includegraphics[width=0.9\\textwidth]{earhart_before_equinox.jpg} \\caption{} \\label{fig:earhart_before} \\end{subfigure}% \\qquad \\begin{subfigure}[t]{0.7\\textwidth} \\centering \\includegraphics[width=0.7\\textwidth]{earhart_near_equinox.jpg} \\caption{} \\label{fig:earhart_equinox} \\end{subfigure}% \\caption{The propeller moonlet Earhart near the Enke gap. (a) This image was taken by the Cassini spacecraft on April 11, 2008 with Cassini's narrow angle camera. The resolution is about $2$ km per pixel. The propeller structure is about $5$ km in radial dimension and about $60$ km in azimuthal dimension. (b) Image of Earhart near Saturn's equinox, taken by Cassini's narrow angle camera on August 13, 2009. The resolution is $7$ km/pixel. The shadow cast by the propeller is $350$km long and the height of the propeller is estimated to be $260$m \\citep{tiscareno10}. Credit: NASA/JPL/Space Science Institute}\\label{fig:earhart} \\end{adjustwidth}% \\end{figure}  In the summer of 2009, near Saturn's equinox, the Cassini spacecraft took images, which show prominent shadows cast by propeller moonlets. The height of the propeller features were calculated from the observed shadow length \\citep{tiscareno10}. For the propeller moonlet Bl\\'{e}riot, a height of $430\\pm30$m was determined, for Santos-Dumont $120$m, and for Earhart $260$m. Figure \\ref{fig:earhart} shows the propeller moonlet Earhart orbiting near the Encke gap, where Figure \\ref{fig:earhart_equinox} portraits Earhart in August 2009, a few days after Saturn's equinox, to cast a $350$km long shadow. In contrast, Figure \\ref{fig:earhart_before} shows Earhart in April 2008 long before equinox.  Using our values of the gap averaged propeller height from Table \\ref{table:heights}, the height of the propeller structure is approximately $0.7$ times the Hill radius of the propeller moonlet. This gives $370$m for the Hill radius of Earhart and $615$m for the Hill radius of Bl\\'eriot. Assuming a sperical shape, the moonlet radius of Earhart is then $240$m for a mass density of $600\\,\\text{kg}/\\text{m}^3$, and $280$m for a mass density of $400\\,\\text{kg}/\\text{m}^3$.  We determined the azimuthal extent of Earhart from Figure \\ref{fig:earhart_equinox} to be $120\\pm20$km, the Cassini ISS team reported about $130$km on their website. The propeller structure thus azimuthally extends about $60\\pm10$km downstream from the moonlet. In Section \\ref{sec:energy_balance} we used a ring temperature of $1.05\\cdot T_\\text{eq}$ to determine the end of the exponential temperature decay. Using this criterium for a moonlet with $370$m Hill radius and for the radial ring particle position $|\\tilde{x}| = 2.5$ (approximately the middle of the gap region), the exponential temperature decay stops after about $7$ orbits. These $7$ orbits at $|\\tilde{x}| = 2.5$ corresponds to an azimuthal propeller length of $61$km downstream of the moonlet which agrees astonishingly well with the $60\\pm10$km determined from Figure \\ref{fig:earhart_equinox}. To calculate the mean decay constant $\\langle \\gamma \\rangle$, we used equation (\\ref{eq:ymax}) with the determined azimuthal propeller extent, to find $\\langle\\gamma\\rangle =  0.05$.  At $|\\tilde{x}| = 2.5$ this corresponds to $k_3 (1-\\varepsilon^2) = 1.1$, which nicely compares to $k_3 (1-\\varepsilon^2) = 1.125$ used in our model."
    },
    "1208/1208.5686_arXiv.txt": {
        "abstract": "Surveys open up unbiased discovery space and generate legacy datasets of long-lasting value. One of the goals of imaging arrays of Cherenkov telescopes like CTA is to survey areas of the sky for faint very high energy gamma-ray (VHE) sources, especially sources that would not have drawn attention were it not for their VHE emission ({\\em e.g.} the Galactic ``dark accelerators\"). More than half the currently known VHE sources are to be found in the Galactic plane. Using standard techniques, CTA can carry out a survey of the region $|\\ell|\\leq 60^\\circ$, $|b|\\leq 2^\\circ$ in 250 hr (1/4th the available time per year at one location) down to a uniform sensitivity of 3 mCrab (a ``Galactic Plane survey\"). CTA could also survey 1/4th of the sky down to a sensitivity of 20 mCrab in 370 hr of observing time (an ``all-sky survey\"), which complements well the surveys by the {\\em Fermi}/LAT at lower energies and extended air shower arrays at higher energies. Observations in (non-standard) divergent pointing mode may shorten the ``all-sky survey'' time to about 100 hr with no loss in survey sensitivity. We present the scientific rationale for these surveys, their place in the multi-wavelength context, their possible impact and their feasibility. We find that the Galactic Plane survey has the potential to detect hundreds of sources. Implementing such a survey should be a major goal of CTA. Additionally, about a dozen blazars, or counterparts to {\\em Fermi}/LAT sources, are expected to be detected by the all-sky survey, whose prime motivation is the search for extragalactic ``dark accelerators\". ",
        "introduction": "Surveys constitute an unbiased, systematic exploratory approach; they favour discoveries of unknown source classes; they allow for scheduling ease and homogeneous data reduction; they provide legacy datasets for future reference. Surveys of different extents and depths are amongst the scientific goals of all major facilities that are planned or in operation. This is particularly critical for observational domains that are opening up, such as very high energy (VHE $\\geq$30 GeV) gamma rays, with wide scope for surprises. Indeed, the Galactic Plane survey carried out by HESS led to the detection of dozens of sources, many of which were unexpected; among these, the {\\em dark accelerators}, have no obvious counterparts at other wavelengths \\citep{Aharonian:2005dx,2006ApJ...636..777A,Aharonian:2008ap}. In high energy gamma rays (HE $\\geq$30 MeV), the {\\em Fermi}/LAT catalog \\citep{2010ApJS..188..405A} has a major impact on our knowledge of the HE sky with statistical studies rendered possible for several classes of sources (blazars, pulsars, globular clusters and normal galaxies), with HE emission associated with unexpected objects (e.g. nova V407 Cyg), with $\\approx 30\\%$ of the 1873 HE sources listed in the second catalog unassociated with known objects \\cite{2011arXiv1108.1435T}.  Compared to previous imaging arrays of Cherenkov telescopes (IACTs), surveys with CTA can only benefit from the increased sensitivity (detection of fainter sources), larger field-of-view (to study multiple or extended sources), improved angular resolution (to alleviate source confusion), broader energy range and better energy resolution (to help determination of the source spectral energy distribution). Surveys provide an immense, if not necessary, service to the research community in the context of an open observatory.  Surveys constitute versatile datasets that enable the detection of unexpected sources and provide testing ground for new theoretical ideas. Surveys are an indispensable tool to assist the community in formulating open time proposals for in-depth studies.  Here, we review current work and perspectives on possible surveys with CTA, their advantages and drawbacks, their relationship with current state-of-the-art and their place in the multi-wavelength context. More precisely, we focus on two easily-defined general purpose surveys that may serve as flagship projects for CTA:  a deep {\\em Galactic Plane survey} and a more shallow, wider {\\em all-sky survey} (both being limited in practice by the fraction of the sky accessible at zenith angle $\\leq 60^\\circ$ from the chosen CTA sites in the Northern and Southern hemispheres). The general scientific objectives and multi-wavelength context are described in \\S2. Simulations have been carried out to study the implementation and achievable sensitivities of these surveys using the latest response files for CTA (\\S3). Their potential in terms of number of detections to expect, based on the current knowledge of various source populations, is then presented in \\S4.  We conclude on the strengths and limitations of both survey proposals. ",
        "conclusions": "CTA will allow a survey of the inner Galactic Plane to unprecedented sensitivity ($\\approx 3$ mCrab), close to the confusion limit, using $\\approx 250$ hr of observing time. Simulations find hundreds of sources can be detected by the survey, enabling population studies and to pinpoint the most interesting sources for deeper follow-up. A Galactic Plane survey should be a major objective of CTA. A wide-area ``all-sky\" survey down to 20 mCrab is also feasible using $\\approx 400$ hr of observing time using standard techniques, or 100 hr using divergent pointing mode. Detailed studies of this mode, which takes advantage of the large number of telescopes in the CTA array, remain to be carried out. The prime motivation for such a survey is the search for new, unsuspected classes of VHE-bright sources (extragalactic ``dark accelerators\") --- admittedly a gamble, but one with large payoff."
    },
    "1208/1208.0540_arXiv.txt": {
        "abstract": "The parsec-scale radio properties of 232 active galactic nuclei (AGNs), most of which are blazars, detected by the Large Area Telescope (LAT) on board the \\textit{Fermi Gamma-ray Space Telescope} have been observed contemporaneously by the Very Long Baseline Array (VLBA) at 5 GHz.  Data from both the first 11 months (1FGL) and the first 2 years (2FGL) of the \\textit{Fermi} mission were used to investigate these sources' $\\gamma$-ray properties.  We use the ratio of the $\\gamma$-ray to radio luminosity as a measure of $\\gamma$-ray loudness.  We investigate the relationship of several radio properties to $\\gamma$-ray loudness and to the synchrotron peak frequency.  There is a tentative correlation between $\\gamma$-ray loudness and synchrotron peak frequency for BL Lac objects in both 1FGL and 2FGL, and for flat-spectrum radio quasars (FSRQs) in 2FGL.  We find that the apparent opening angle tentatively correlates with $\\gamma$-ray loudness for FSRQs, but only when we use the 2FGL data.  We also find that the total VLBA flux density correlates with the synchrotron peak frequency for BL Lac objects and FSRQs.  The core brightness temperature also correlates with synchrotron peak frequency, but only for the BL Lac objects.  The low-synchrotron peaked (LSP) BL Lac object sample shows indications of contamination by FSRQs which happen to have undetectable emission lines.  There is evidence that the LSP BL Lac objects are more strongly beamed than the rest of the BL Lac object population. ",
        "introduction": "The Large Area Telescope (LAT; Atwood et al.\\ 2009) on board the \\textit{Fermi Gamma-ray Space Telescope} is a wide-field telescope covering the energy range from about 20 MeV to more than 300 GeV.  It has been scanning the entire $\\gamma$-ray sky once every three hours since July of 2008, with breaks for flares and other targets of opportunity.  The majority of the sources (685 of 1451) in the \\textit{Fermi} LAT First Source Catalog (1FGL; Abdo et al. 2010a) have been identified with known radio blazars.  These blazars typically are active galactic nuclei (AGNs) with strong, compact radio sources which exhibit flat radio spectra, rapid variability, compact cores with one-sided parsec-scale jets, and superluminal motion in the jets (Marscher 2006).  This trend continues in the newly released \\textit{Fermi} LAT Second Source Catalog (2FGL; Nolan et al. 2011), compiled using the first two years of LAT data.  We previously presented findings on the relationships between the LAT-detected and non-LAT-detected populations of blazars (Linford et al. 2011; Linford et al. 2012).  Like other studies, we found a strong correlation between LAT flux and total VLBA radio flux density.  We also found that the LAT and non-LAT BL Lac objects appeared to be similar in many respects, while the LAT flat-spectrum radio quasars (FSRQs) were extreme sources when compared to their non-LAT counterparts.  Polarized emission at the base of the jets was also reported to be significantly more frequent in LAT blazars than in non-LAT blazars.  A major program to monitor the parsec-scale radio (15 GHz) properties of these $\\gamma$-ray emitting blazars is the Monitoring Of Jets in AGN with VLBA Experiments (MOJAVE; Lister et al.\\ 2009a; Homan et al.\\ 2009).  The MOJAVE and \\textit{Fermi} LAT collaborations recently published a paper detailing their investigations of parsec-scale properties of the $\\gamma$-ray emitting blazars in their sample (Lister et al. 2011).  They studied the relationships between radio properties and the $\\gamma$-ray loudness ($G_{r}$; the ratio of $\\gamma$-ray-to-radio luminosity) and synchrotron peak frequency ($\\nu^{S}_{peak}$; the frequency where the synchrotron emission is at maximum).  They reported significant differences in the $G_{r}$ distributions between the BL Lac objects and FSRQs.  For their BL Lac objects, they reported strong correlations for $G_{r}$-$\\nu^{S}_{peak}$ and $G_{r}$-$\\gamma$-ray photon spectral index.  They also reported a non-linear correlation between apparent jet opening angle and $G_{r}$ for their entire sample.  Their high-synchrotron peaked ($\\nu^{S}_{peak} > 10^{15}$) BL Lac objects tended to have lower core brightness temperatures, linear core polarization, and variability than the rest of their BL Lac object population.  Here, we have analyzed our recent VLBA 5 GHz data, taken contemporaneously with LAT observations (Linford et al. 2012), to see if we could reproduce the findings of Lister et al. (2011).  Our sample was slightly larger than the one presented in Lister et al. (2011), with 232 sources in our sample compared to 173 in theirs.  We had a larger fraction of BL Lac objects in our sample, and we also had more non-blazar AGNs (radio galaxies, AGN of unknown type, and one starburst galaxy).  The objects in our sample spanned a larger range of radio flux densities, as the MOJAVE sample is targeting the brightest AGN and our sample is radio flux limited.  We only had single observations on our sources, as we are not a monitoring program.  This leads to some difficulties in comparing with the MOJAVE sample, especially in terms of apparent jet opening angle.  Another area that has garnered renewed interest in the \\textit{Fermi}-LAT era is the unification system for AGN (e.g., Urry \\& Padovani 1995).  Several groups (e.g., Nieppola et al. 2008, Ghisellini \\& Tavecchio 2008, Meyer et al. 2011, and Giommi et al. 2012b) have been investigating the ``blazar sequence'' (Fossati et al. 1998, Ghisellini et al. 1998), the relationship between blazar luminosity and synchrotron peak frequency.  There have been hints (Vermeulen et al. 1995, Ghisellini et al. 2009, Giommi et al. 2012a, Meyer et al. 2011) that the population of BL Lac objects with low (below 10$^{14}$ Hz) synchrotron peak frequencies might actually contain some misidentified FSRQs which happen to have strong emission from their jets overpowering their emission lines.  We investigated the possibility for this kind of ``contamination'' in our sample and present our findings here.  In Section 2 we define our sample.  In Section 3 we discuss how we determined the $\\gamma$-ray loudness and synchrotron peak frequency for our sources. In Section 4 we present our results and discuss their implications.  In Section 5, we present evidence that our LSP BL Lac object sample may have some FSRQs hiding in it.  Throughout this paper we assume $\\Lambda$CDM cosmology with $H_{0} = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m} = 0.27$, and $\\Omega_{\\Lambda} = 0.73$ (e.g., Hinshaw et al.\\ 2009). ",
        "conclusions": "We have analyzed a sample of 232 LAT-detected AGN using both 1FGL and 2FGL data to compare our results with those of Lister et al. (2011).  All of the sources in our sample are significantly $\\gamma$-ray loud.  We did not find a significant difference between the distributions of $G_{r}$ for BL Lac objects and FSRQs.  Using 1FGL data, we find a weak correlation between $G_{r}$ and $\\nu^{S}_{peak}$ for the BL Lac objects and a tentative correlation for the FSRQs.  Using 2FGL, we found a very tentative $G_{r}$-$\\nu^{S}_{peak}$ correlation for both FSRQs and BL Lac objects.  Looking at the parsec-scale radio properties of our sources, we find a very strong correlation between total VLBA flux density and $\\nu^{S}_{peak}$ for both BL Lac objects and FSRQs.  We could not confirm the correlation between $G_{r}$ and the $\\gamma$-ray photon index reported by Lister et al. (2011), but we did confirm the correlation between $\\gamma$-ray luminosity and $\\gamma$-ray photon index reported by other groups.  We also found a very strong correlation between the core brightness temperatures and $\\nu^{S}_{peak}$ for BL Lac objects.  Although we had a limited sample of apparent jet opening angle measurements, we were still able to tentatively confirm the correlation with $G_{r}$ reported by Lister et al. (2011).  We did not find any evidence of a correlation between core fractional polarization and $\\nu^{S}_{peak}$.  We found a tentative negative correlation between radio variability (modulation index) and $\\nu^{S}_{peak}$ for the BL Lac objects and a tentative positive correlation for the FSRQs.  The fact that the core brightness temperature shows a positive correlation with $\\nu^{S}_{peak}$ and modulation index shows a negative correlation with $\\nu^{S}_{peak}$ indicates that the LBLs are more strongly beamed than the IBLs and HBLs.  The LBLs in our sample often appear to be different for the rest of the BL Lac objects.  In particular, we found significant differences in the distributions of core brightness temperatures and total VLBA flux density.  % It seems likely, therefore, that our LBL population contains some misidentified FSRQs which may have their BLR swamped by their jet emission.  While Lister et al. (2011) argued that a $G_{r}$-$\\nu^{S}_{peak}$ correlation for the BL Lac objects indicated that the LBLs were related to the IBLs and HBLs, we found that deliberately contaminating our LBL sample with known FSRQs did not change our (albeit weak) correlation significantly.  Future studies of large samples of blazars, which should include both very high and low flux density objects, should be conducted to further investigate the relationships between $G_{r}$, $\\nu^{S}_{peak}$, and the parsec-scale radio properties.  Long-term monitoring of LBLs may also present clear evidence that some of these objects are actually FSRQs.  \\noindent We thank the anonymous referee for their constructive criticism and helpful comments on this manuscript. We thank Talvikki Hovatta for useful discussions regarding MOJAVE core fractional polarization measurements and Roger Romani for useful discussions about the 2LAC synchrotron peak frequency estimation method. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. The VLBA is a facility of the National Science Foundation operated by the National Radio Astronomy Observatory under cooperative agreement by Associated Universities, Inc. MATLAB is a registered trademark of The MathWorks, Inc. (Natick, Massachusetts, USA). The NASA/IPAC Extragalactic Database (NED) is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We thank NASA for support under FERMI grant GSFC \\#21078/FERMI08-0051, and the NRAO for support under Student Observing Support Award GSSP10-011.  Additional support provided by the Naval Research Laboratory."
    },
    "1208/1208.5709_arXiv.txt": {
        "abstract": "We present a mass determination for the transiting super-Earth $\\rho^{1}$~Cancri~e based on nearly 700 precise radial velocity (RV) measurements. This extensive RV data set consists of data collected by the McDonald Observatory planet search and published data from Lick and Keck observatories (Fischer et al.~2008). We obtained 212 RV measurements with the Tull Coud\\'e Spectrograph at the Harlan J. Smith 2.7\\,m Telescope and combined them with a new Doppler reduction of the 131 spectra that we have taken in 2003-2004 with the High-Resolution-Spectrograph (HRS) at the Hobby-Eberly Telescope (HET) for the original discovery of $\\rho^{1}$~Cancri~e. Using this large data set we obtain a 5-planet Keplerian orbital solution for the system and measure an RV semi-amplitude of $K=6.29\\pm0.21$\\,m\\,s$^{-1}$ for $\\rho^{1}$~Cnc~e and determine a mass of 8.37$\\pm0.38$~M$_{\\oplus}$. The uncertainty in mass is thus less than $5\\%$. This planet was previously found to transit its parent star (Winn et al.~2011, Demory et al.~2011), which allowed them to estimate its radius. Combined with the latest radius estimate from Gillon et al.~(2012), we obtain a mean density of $\\rho=4.50\\pm0.20$\\,g\\,cm$^{-3}$. The location of $\\rho^{1}$~Cnc~e in the mass-radius diagram suggests that the planet contains a significant amount of volitales, possibly a water-rich envelope surrounding a rocky core. ",
        "introduction": "The $\\rho^{1}$~Cancri planetary system is one of the most interesting, nearby multi-planet systems that was discovered and extensively studied by the radial velocity (RV) technique. The parent star, $\\rho^{1}$~Cancri (= HR~3522, HD~75732, 55~Cancri), is a $V=5.95$ G8V (Montes et al.~2001) star, located at a distance of $12.3\\pm0.1$\\,pc, based on the Hipparcos parallax of $81.03\\pm0.75$~mas (Van Leeuwen~2007). It is the primary of a wide visual binary with a projected separation of $\\approx1065$~AU (Mugrauer et al. 2006). The first planet in this system ($\\rho^{1}$~Cnc~b, $P=14.65$~d) was found by Butler et al.~(1996), based on Lick Observatory RV data. Six years later, Marcy et al.~(2002) presented evidence for two more giant planets orbiting this star: $\\rho^{1}$~Cnc~c with a period of 44~d and a very long period planet at an orbital separation of 5.5~AU, $\\rho^{1}$~Cnc~d. An intensive RV campaign that we carried out using the Hobby-Eberly Telescope (HET), revealed a short-periodic signal that was, at that time, thought to be one of the first discoveries of a hot Neptune with a minimum mass of $\\sim 17$~M$_{\\oplus}$ and an orbital period of 2.8 days (McArthur et al. 2004, hereafter Mc04), raising the total number of detected planets in this system to four. The planet count was further increased by Fischer et al.~(2008, hereafter F08), who presented evidence for a fifth planet with an orbital period of 260~days.  Dawson \\& Fabrycky (2010, hereafter DF10) re-analyzed the published RV time-series data for this star and claimed that the 2.8~day period of $\\rho^{1}$~Cnc~e is an alias and that the true period of this companion is just 0.7365 days. This shorter period also led to a reduction of the minimum mass of $\\rho^{1}$~Cnc~e to around $8$~M$_{\\oplus}$, moving the planet from the Neptune mass range into the super-Earth mass bin. DF10 pointed out the high a-priori transit probability of 25\\% for a planet and motivated highly precise photometric observations to search for the planetary transit signal. This transit search was indeed successful, using two different space telescopes, MOST and warm {\\it Spitzer}. Winn et al.~(2011, hereafter W11), and Demory et al.~(2011, hereafter D11) practically simultaneously announced the detection of the transit signal of this super-Earth. The transit thus confirmed the shorter orbital period for $\\rho^{1}$~Cnc~e, as suggested by DF10. D11 measure a planetary radius of $2.1\\pm0.17~R_{\\oplus}$ and a mean density of $4.8\\pm1.3$\\,g\\,cm$^{-3}$, while W11 reported a similar radius of $2.0\\pm0.14~R_{\\oplus}$ and slightly higher mean density of $5.9^{+1.5}_{-1.1}$\\,g\\,cm$^{-3}$. This places $\\rho^{1}$~Cnc~e in the mass-radius diagram between the region of high-density, rocky planets like CoRoT-7b (Hatzes et al. 2011) and Kepler-10b (Batalha et al.~2011) and planets with a significant amount of volatiles, so-called ``Mini-Neptunes'', like GJ~1214b (Charbonneau et al.~2009) and the Kepler-11 planets (Lissauer et al.~2011). An improved radius determination of $2.17\\pm0.10~R_{\\oplus}$ was presented by Gillon et al.~(2012, hereafter G12) by combining the {\\it Spitzer} with the MOST photometry. Even more recently, Demory et al.~(2012) used Warm {\\it Spitzer} 4.5~$\\mu$m observations of occultations of $\\rho^{1}$~Cnc~e to detect its thermal emission. Clearly, this nearby transiting super-Earth planet around a bright star offers an abundant variety of very interesting follow-up observations that will allow a detailed characterization of this exoplanet.  In our paper we will focus on the mass determination for $\\rho^{1}$~Cnc~e based on hundreds of precise RV measurements. The paper is structured as follows: we first describe the observations of $\\rho^{1}$~Cnc at McDonald Observatory, the second section contains a description of the multi-planet orbital fit that we have performed, and the third section discusses our results, in particular the precise mass, for $\\rho^{1}$~Cnc~e. ",
        "conclusions": "With the inclination $i$ of 82.5 degrees measured by G12 from the photometric transit, we derive a planetary mass of $8.37\\pm0.38~M_{\\oplus}$ for $\\rho^{1}$~Cnc~e. This value is slightly higher than reported by D11 ($7.81\\pm0.56~M_{\\oplus}$) and slightly lower than the value of W11 ($8.63\\pm0.35~M_{\\oplus}$). The W11 mass and its error were taken from the orbital solutions of DF10. With our approach, we could not reproduce the small error in $K$ of only 0.2\\,m\\,s$^{-1}$ presented by DF10 using the Keck and Lick RV data alone (we obtain an uncertainty in $K$ of 0.33\\,m\\,s$^{-1}$).  Our mass estimate is based on twice the amount of RV data and we also benefit from the improved stellar mass determination of von Braun et al.~(2011). We therefore regard this new mass determination as the current best value for $\\rho^{1}$~Cnc~e. The uncertainty in mass is driven by the uncertainty in $K$ and in second order coupled to the uncertainty in the mass of the host star. We are thus only limited by the precision of the radial velocity measurements and the quality of the stellar parameters of $\\rho^{1}$~Cnc.  Figure\\,\\ref{mr} shows $\\rho^{1}$~Cnc~e in the mass-radius diagram compared to models for internal composition of small and low-mass planets from Seager et al.~(2007) and three other transiting planets with well determined masses and radii. $\\rho^{1}$~Cnc~e has the smallest area of uncertainty based on its errors in mass and radius (Kepler-10b has a smaller radius error but a larger uncertainty in mass). As noted by other authors (e.g. G11), $\\rho^{1}$~Cnc~e requires a significant amount of volatiles to explain its location in this diagram. It is located significantly above the model curves for purely rocky planets and approaches the zone of ``mini-Neptunes''. Still, it does not require a large H/He envelope, but its mass and radius rather suggest a water-rich envelope around a rocky core.  Kaib, Raymond \\& Duncan~(2011) suggest that the $\\rho^{1}$~Cnc planetary system is coplanar but misaligned with its host star spin axis due to the perturbations of the secondary star. In principle, this can be tested (at least for $\\rho^{1}$~Cnc~e) by observing and measuring the Rossiter-McLaughlin (RM) effect. However, the expected amplitude of the RM effect for $\\rho^{1}$~Cnc~e is only 0.5\\,m\\,s$^{-1}$, and while RV measurements were obtained during transit by coincidence, a signal of this small amplitude is clearly undetectable by the current data. Future extreme precision RV measurements of several transits (by e.g. the upgraded HET/HRS or HARPS-North) might allow us to measure the spin-orbit misalignment for this planet.  As mentioned before, despite the large quantity of precise RV measurements we did not detect any significant residual signal that could indicate a sixth planet in the system. The most interesting peak is near 131 days, as the window function is clean at this period value, and it would be at an orbital separation between the 43~d and the 261~d planet. But this peak has only a modest power and is not statistically signifcant. However, there is a rapidly increasing ``treasure trove'' of precise RV measurements for this system, with our paper adding over 300 RV points to the published sample. This should allow in the future to achieve sensitivity for more planets, either with lower mass or at longer orbital periods and especially in the habitable zone and in the currently large empty region between the inner 4 planets and the distant outer planet at $a\\approx5$~AU. The $\\rho^{1}$~Cnc multi-planetary system likely has more exciting discoveries waiting to be made."
    },
    "1208/1208.4539_arXiv.txt": {
        "abstract": "We report quadrature observations of an extreme-ultraviolet (EUV) wave event on 2011 January 27 obtained by the Extreme Ultraviolet Imager (EUVI) onboard \\emph{Solar Terrestrial Relations Observatory} (\\emph{STEREO}), and the Atmospheric Imaging Assembly (AIA) onboard the \\emph{Solar Dynamics Observatory} (\\emph{SDO}). Two components are revealed in the EUV wave event. A primary front is launched with an initial speed of $\\sim$440 km s$^{-1}$. It appears significant emission enhancement in the hotter channel but deep emission reduction in the cooler channel. When the primary front encounters a large coronal loop system and slows down, a secondary much fainter front emanates from the primary front with a relatively higher starting speed of $\\sim$550 km s$^{-1}$. Afterwards the two fronts propagate independently with increasing separation. The primary front finally stops at a magnetic separatrix, while the secondary front travels farther before it fades out. In addition, upon the arrival of the secondary front, transverse oscillations of a prominence are triggered. We suggest that the two components are of different natures. The primary front belongs to a non-wave coronal mass ejection (CME) component, which can be reasonably explained with the field-line stretching model. The multi-temperature behavior may be caused by considerable heating due to the nonlinear adiabatic compression on the CME frontal loop. For the secondary front, most probably it is a linear fast-mode magnetohydrodynamic (MHD) wave that propagates through a medium of the typical coronal temperature. X-ray and radio data provide us with complementary evidence in support of the above scenario. ",
        "introduction": "One of the most intriguing phenomena discovered by the Extreme-ultraviolet (EUV) Imaging Telescope \\citep[EIT;][]{Delab95} onboard the \\emph{Solar and Heliospheric Observatory} (\\emph{SOHO}) satellite is ``EIT waves\", which are characterized by a diffuse bright front globally propagating through the solar corona \\citep{Moses97,Thompson98}.  EIT waves were initially interpreted as a fast-mode magnetohydrodynamic (MHD) wave in the corona \\citep{Thompson99}, which can travel across the magnetic field lines freely, covering a quite large fraction of the solar disk. If the coronal fast-mode wave is strong enough, it can also perturb the much denser chromosphere at its base to produce an H$\\alpha$ Moreton wave, just as the scenario proposed by \\citet{Uchida68}. Many subsequent numerical and observational studies \\citep[e.g.,][]{Wang00,Wu01,Warmuth04,Veronig06,Long08,Gopalswamy09,Patsourakos09a} have provided further evidence for this view.  Such a fast-mode wave model was first challenged by \\citet{Delannee99} who found that an EIT wave stopped at the magnetic separatrix, which is hard to explain in the wave framework. In addition, case studies have revealed that the EIT wave front is co-spatial with the coronal mass ejection (CME) frontal loop \\citep[e.g.,][]{Attrill09,Chen09,Dai10}. Hence several alternative models have been proposed, which regard EIT waves as a result of magnetic reconfiguration related to the CME liftoff rather than a true wave in the corona. These non-wave models include the current shell model \\citep{Delannee00}, the field-line stretching model \\citep{Chen02,Chen05}, and the successive reconnection model \\citep{Attrill07}. Besides, some other authors claim EIT waves to be a type of slow-mode MHD wave \\citep{Wills-Davey07,Wang09}. For more details on the observations and modeling of EIT waves, please refer to recent reviews \\citep{Wills-Davey09,Gallagher11,Zhukov11,Chen11b,Patsourakos12}.  \\citet{Chen02} predicted that there should be a fast-mode wave ahead of the EIT wave, which was confirmed by \\citet{Harra03}. On the other hand, \\citet{Zhukov04} suggested from the observational point of view that there could be both wave and non-wave components in an EIT wave. However, early EIT wave studies to catch such multiple components often suffered the low cadence of EIT, which is 12 minutes at best. The situation has been greatly improved with the launch of the \\emph{Solar Terrestrial Relations Observatory} \\citep[\\emph{STEREO};][]{Kaiser08} and the \\emph{Solar Dynamics Observatory} \\citep[\\emph{SDO};][]{Pesnell12}. Thanks to the much higher temporal resolutions of the EUV telescopes onboard the three spacecraft, multiple components in an EIT wave have been successfully identified in observations \\citep[e.g.,][]{Liu10,Chen11a,Cheng12,Asai12} and verified in numerical efforts \\citep[e.g.,][]{Cohen09,Downs11,Downs12}. With the observations of modern generation of EUV imagers, now we prefer the more general term ``EUV wave\" to the conventional one ``EIT wave\". In this paper we report quadrature observations of two components and their decoupling in an EUV wave event on 2011 January 27 from both \\emph{STEREO} and \\emph{SDO}\\@. The distinct differences in amplitude, kinematics, and multi-temperature behavior imply their different physical mechanisms. In Section 2 we introduce the instruments and data sets. Analysis is carried out and results are presented in Section 3. Then we discuss the results in Section 4 and draw our conclusions in Section 5. ",
        "conclusions": "By using the \\emph{STEREO-A}/EUVI and \\emph{SDO}/AIA quadrature observations of an EUV wave event on 2011 January 27, two fronts and their decoupling are revealed. The two fronts show distinct differences in amplitude, kinematics, and multi-temperature behavior. Complemented with the X-ray and radio observations, we suggest that the two fronts are of different natures. The primary front belongs to a non-wave CME component, which can be reasonably explained with the field-line stretching model. For the secondary front, most probably it is a linear fast-mode MHD wave that propagates through a medium of the typical coronal temperature. The decoupling of the two fronts is caused by the interaction of the CME frontal loop and a large coronal loop system south of it."
    },
    "1208/1208.4363_arXiv.txt": {
        "abstract": "We present stellar mass surface density profiles of a mass-selected sample of 177 galaxies at $0.5 < z < 2.5$, obtained using very deep HST optical and near-infrared data over the GOODS-South field, including recent CANDELS data. Accurate stellar mass surface density profiles have been measured for the first time for a complete sample of high-redshift galaxies more massive than $10^{10.7} M_\\odot$.  The key advantage of this study compared to previous work is that the surface brightness profiles are deconvolved for PSF smoothing, allowing accurate measurements of the structure of the galaxies. The surface brightness profiles account for contributions from complex galaxy structures such as rings and faint outer disks. Mass profiles are derived using radial rest-frame $u\u2212g$ color profiles and a well-established empirical relation between these colors and the stellar mass-to-light ratio. We derive stellar half-mass radii from the mass profiles, and find that these are on average $\\sim25\\%$ smaller than rest-frame $g$ band half-light radii. This average size difference of 25\\% is the same at all redshifts, and does not correlate with stellar mass, specific star formation rate, effective surface density, S\\'ersic index, or galaxy size.  Although on average the difference between half-mass size and half-light size is modest, for approximately 10\\% of massive galaxies this difference is more than a factor two. These extreme galaxies are mostly extended, disk-like systems with large central bulges. These results are robust, but could be impacted if the central dust extinction becomes high. ALMA observations can be used to explore this possibility. These results provide added support for galaxy growth scenarios wherein massive galaxies at these epochs grow by accretion onto their outer regions. ",
        "introduction": "Over the past decades quantitative studies of high-redshift galaxy structure have advanced tremendously. Sensitive, high resolution instruments such as the Hubble Space Telescope's (\\textit{HST}) Advanced Camera for Surveys (ACS) and the Wide Field Camera 3 (WFC3) have made it possible to measure the structure of faint high-redshift galaxies at sub-kpc scales. Furthermore, the availability of easy-to-use photometric redshift and stellar population fitting packages has made it possible to straightforwardly measure a wide variety of parameters for ever-increasing numbers of galaxies.  Since a small amount of recent star formation can have a disproportionately large contribution to a galaxy's light compared to its mass, galaxies are usually observed at the redmost wavelengths, where emission from young stars is weakest. At low redshift, this can be done quite effectively, since rest-frame near-infrared (NIR) data is available at high enough resolution. At higher redshifts, however, it is impossible to observe at such long wavelengths with sufficiently high angular resolution. Until recently, the Hubble Space Telescope (\\textit{HST}) Advanced Camera for Surveys (ACS) was the only wide-field instrument capable of measuring the structure of high-redshift galaxies in any detail. At $z=2$, the reddest filter available on ACS, $Z_{850}$, corresponds to rest-frame near-ultraviolet (NUV) wavelengths. Use of such short-wavelength data has been shown to result in drastically different conclusions about galaxy structure and morphology, compared to rest-frame optical data (e.g., \\citealt{lab03}; \\citealt{tof05}; \\citealt{cam11}).  With the introduction of the \\textit{HST} Wide Field Camera 3 (WFC3) it has become possible to measure rest-frame optical light of $z\\sim2$ galaxies at a resolution approaching that of the ACS. The redder light detected by this instrument provides a much better proxy for stellar mass. However, color gradients are known to exist to some extent in all types of galaxies at redshifts up to at least $z\\sim 3$, such that most galaxies contain a relatively red core and blue outer regions (e.g., \\citealt{dok10}; \\citealt{szo11}; \\citealt{guo11}). These color variations are caused by a combination of varying dust content, metallicity and stellar age, and imply that the stellar mass-to-light ratio ($M_*/L$) of a given galaxy varies with position within that galaxy. Thus, even though rest-frame optical light is a better tracer of stellar mass than rest-frame NUV light, neither accounts for the complexity of stellar population variations within galaxies.  By fitting stellar population models to resolved galaxy photometry, it is in principle possible to infer spatial variations in stellar mass, age, metallicity, dust content, and other parameters. This approach is currently somewhat limited by the lack of high-resolution data at infrared wavelengths, but can nonetheless be used to measure several basic stellar population properties. An example of this technique is presented in \\cite{wuy12}, who have performed stellar population modeling on resolved \\textit{HST} data, using integrated IR observations as constraints on the overall properties of their galaxies. In this approach the integrated photometry serves as an important tool to constrain the overall spectral energy distribution (SED) of a galaxy, while the \\textit{HST} data provide information regarding the spatial variation of the stellar populations within these constraints.  In this Paper we explore an alternative method to recover $M_*/L$ variations, using a simple empirical relation between rest-frame $u-g$ color and $M_*/L$. Using this method we construct stellar mass surface density profiles corrected for the effects of the PSF, for a mass-selected sample of galaxies between $z=0$ and $z=2.5$. We compare the resulting half-mass radii to half-light radii based on rest-frame optical imaging. All sizes presented in this Paper are circularized sizes: $r_e = r_{e,a}\\sqrt{b/a}$. Throughout the Paper we assume a $\\Lambda$CDM cosmology with $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We have presented the first consistently measured stellar mass surface density profiles for individual galaxies at redshifts between $z=0$ and $z=2.5$. These profiles have been derived using an empirical relation between rest-frame color and stellar mass-to-light ratio. This simple method does not yield the same detailed information regarding, e.g., stellar ages and dust content as resolved SED-fitting techniques. However, it is robust to variations in stellar population properties; changes in stellar age, metallicity, or dust content shift galaxies roughly along the empirical relation, and are thus implicitly included in our $M_*/L$ determination.  The key advantage of this study compared to previous work at high redshift (e.g., \\citealt{wuy12}; \\citealt{lan12}) is the fact that the profiles presented here are deconvolved for PSF smoothing. This is crucial for measurements of high-redshift galaxy structure, since these distant and often physically small galaxies have angular sizes that are in many cases comparable to the \\textit{HST} PSF size (e.g., \\citealt{dad05}; \\citealt{tru06}; \\citealt{tof07}; \\citealt{dok08}; \\citealt{szo10}; \\citealt{cas10}; \\citealt{szo12}). Since our surface density profiles are derived from deconvolved surface brightness profiles, they can be used to correctly measure structural parameters such as sizes and S\\'ersic indices.  The considerable depth of the data used in this study allows us to probe galaxy fluxes and colors out to large radii. The robustness of the resulting structural parameters has been tested using ultradeep data taken over the HUDF, which overlaps with the CANDELS GOODS-South field. A comparison of galaxy parameters derived using the two datasets confirms that our results are not systematically affected by surface brightness effects.  We have shown that the half-mass radii of galaxies between $z=0$ and $z=2.5$ are on average 25\\% smaller than their rest-frame optical half-light radii. There is significant scatter in this size difference, with some galaxies having half-mass radii that are almost an order of magnitude smaller than their half-light radii. We find that, on average, this size difference does not vary with redshift for galaxies at fixed mass. This holds for the population as a whole, as well as for the quiescent and starforming subpopulations seperately. This is an interesting result, as it implies that $z\\sim2$ galaxies have similar color gradients as their low-redshift analogs, despite the fact that these low-redshift galaxies formed at a different epoch, and perhaps through very different formation mechanisms.  There does not seem to be a strong correlation between galaxy morphology or star forming activity and the difference between half-mass size and half-light size. However, we do find that the galaxies with the most extreme size differences are almost all extended disk galaxies with very prominent central bulges. These galaxies range from strongly starforming to almost completely quiescent, and may represent a short transitional phase during which the central bulge is prominent and the starforming disk is very young.  There is significant scatter around the empirical relation used to convert $u-g$ color to $M_*/L$, which could lead to a systematic underestimate of the mass-to-light ratios in galaxy regions that are relatively metal-rich or old. However, this effect is likely small for moderate stellar population variations. Similarly, very high central dust concentrations may result in an underestimate of the inner mass content of galaxies, biasing our results towards larger half-mass radii. The high-resolution infrared data needed in order to quantify such dust effects will be available in the (near) future, with instruments such as ALMA.  Inside-out galaxy growth, as described by, e.g., \\cite{dok10}, implies that the growth of the most massive galaxies since $z\\sim2$ is largely due to material being accreted onto the outer regions of these galaxies. The cores of massive galaxies likely formed in short, violent bursts at higher redshift, and should therefore have star formation histories and stellar populations that are quite different from those in the outer regions. The results presented in this Paper broadly agree with such a picture; the central regions of massive galaxies are redder, and therefore likely older, than the outer regions. Using the method presented in this Paper we cannot, however, disentangle dust, age and metallicity gradients; nor can we constrain the star formation histories within our galaxies. First steps towards a better understanding of stellar population variations within high-redshift galaxies have been made by several authors. Results based on photometry of early-type galaxies (e.g., \\citealt{guo11}; \\citealt{gar12}) indicate that stellar age and metallicity are the dominant drivers of radial color gradients. Studies based on spectroscopic measurements of gravitationally lensed high-redshift galaxies (e.g., \\citealt{cre10}; \\citealt{jon10}; \\citealt{yua11}; \\citealt{que12}; \\citealt{jon12}) have shown that most of these galaxies have negative metallicity gradients. These results seem to be roughly consistent with each other, but are based on very differently selected, and rather small, galaxy samples. A broader, more in-depth analysis of radial stellar population variations, for a well-defined sample of starforming galaxies as well as quiescent galaxies, could provide valuable insights into the processes which have shaped the structure of galaxies today."
    },
    "1208/1208.4377_arXiv.txt": {
        "abstract": "Doppler surveys have shown that the occurrence rate of Jupiter-mass planets appears to increase as a function of stellar mass.  However, this result depends on the ability to accurately measure the masses of evolved stars.  Recently, Lloyd (2011) called into question the masses of subgiant stars targeted by Doppler surveys. Lloyd argues that very few observable subgiants have masses greater than 1.5\\msun, and that most of them have masses in the range 1.0-1.2~\\msun. To investigate this claim, we use Galactic stellar population models to generate an all-sky distribution of stars. We incorporate the effects that make massive subgiants less numerous, such as the initial mass function and differences in stellar evolution timescales. We find that these effects lead to negligibly small systematic errors in stellar mass estimates, in contrast to the $\\approx50$\\% errors predicted by Lloyd. Additionally, our simulated target sample does in fact include a significant fraction of stars with masses greater than 1.5~\\msun, primarily because the inclusion of an apparent magnitude limit results in a Malmquist-like bias toward more massive stars, in contrast to the volume-limited simulations of Lloyd. The magnitude limit shifts the mean of our simulated distribution toward higher masses and results in a relatively smaller number of evolved stars with masses in the range 1.0--1.2~\\msun. We conclude that, within the context of our present-day understanding of stellar structure and evolution, many of the subgiants observed in Doppler surveys are indeed as massive as main-sequence A stars. ",
        "introduction": "Studies of the relationships between exoplanets and their host stars provide valuable clues about how planets form, and also point the way to new discoveries. For example, the well-established relationship between the occurrence rate of gas giant planets and host-star metallicity \\citep{santos04,fischer05b,johnson09b} may be an indication that the formation  timescale for close-in giant planets ($a < 5$~AU) is shortened by the metal-enhancement, and hence dust-enhancement, of protoplanetary disks \\citep[e.g.][]{ida04}. For this reason, certain Doppler surveys have biased their target lists toward metal-rich stars, which has resulted in the discovery of many of the known hot Jupiter systems \\citep{fischer05a, bouchy05, sato05}.  More recent Doppler surveys have discovered that stellar mass is another key predictor of giant planet occurrence \\citep{johnson07b,johnson10c}. This relationship is based on Doppler surveys of M dwarfs on one side of the stellar mass range \\citep[e.g.][]{johnson10a}, and the evolved   counterparts of F- and A-type stars on the more massive end \\citep{johnson07,lovis07,sato07}. These so-called ``retired A-stars'' exhibit dramatically slower rotation velocities (\\vsini) than their main-sequence progenitors \\citep{gray85,donascimento00}, making them better targets for Doppler-based planet surveys compared to their F- and A-type main-sequence counterparts \\citep{hatzes03,fischer03,galland05}.  However, the mass estimates of subgiants targeted by Doppler surveys have recently been called into question by \\citet[][hereafter L11]{lloyd11}. In an attempt to study the effects of star-planet tidal interactions in planetary systems with evolved host stars, L11 investigated the expected mass distribution of evolved stars near the subgiant branch. By using stellar evolution model grids, assumptions about the metallicity distribution in the Galaxy, and the form of the stellar initial mass function (IMF), L11 concluded that most bright subgiants are not the evolved brethren of A-type stars, but rather the evolved counterparts of Sun-like stars. This is because massive stars evolve much more quickly along the subgiant branch than do less massive stars. As L11 notes, this differential evolution rate for stars of different masses is a robust feature of stellar models. L11 predict that this effect, together with the distribution of stellar masses produced by the initial mass function, should  result in a very small number of  massive subgiants with $M \\gtrsim 1.5$~\\msun\\ in Doppler surveys. We note that while L11 discuss stellar rotation in great detail, it is this evolution rate feature that is his key argument that subgiant masses are incorrect. We therefore focus our investigation on this effect with the goal of assessing the question: could the mass estimates of subgiants be systematically overestimated by ignoring the stellar IMF and mass-dependent evolution rate along the subgiant branch?  In this contribution we assess the specific critique of L11 using a simple application of a Bayesian framework to the Galactic population models of \\citet{girardi05}.  We show that the neglect of the IMF and the mass-dependent evolution timescales of subgiants results in a small bias in the mass measurement towards higher masses. But that this bias is too small to cast doubt on the conclusions of \\citet{johnson10b, johnson10c}, namely that the occurrence of Jovian planets increases with increasing stellar mass. We also demonstrate that the mass distribution of the stars in the Johnson et al. Keck Doppler survey is expected to contain a substantial number of subgiants, consistent with the masses measured for that survey and strongly inconsistent with the mass distribution predicted for it by L11. ",
        "conclusions": "\\label{sec:results}  L11 argued that the mass measurements of subgiants with $M > 1.5$~\\msun, i.e. the retired A stars surveyed by \\citet{johnson10b}, must be in error because stars in this mass range should be exceedingly rare. L11 further argues that stellar evolutionary models are sufficiently ambiguous in their predictions (given reasonable uncertainties in their input physics) and that spectroscopically determined stellar parameters are subject to large systematic errors. L11 concludes that the true masses of the stars in the Johnson et al. sample are more reasonably estimated to be 1.0-1.2 solar masses, not typically closer to 1.5 solar masses.  We applied the survey selection criteria of Johnson et al. to the TRILEGAL galactic synthesis models and have shown that the resulting simulated target sample has a mass distribution consistent with the Johnson et al. mass measurements and inconsistent with the prediction of L11.  L11 may be correct that stellar evolution and Galactic synthesis models have substantial uncertainties. However, since the TRILEGAL models successfully and accurately reproduces the stellar characteristics of stars in the Solar Neighborhood \\citep{girardi05}, we find no reason to doubt their accuracy at the level that would implicate the Johnson et al. mass measurements.  Nevertheless, tests of systematic errors in stellar evolution models using planet transit light curves, eclipsing binaries and asteroseismology are very much worthwhile. Fortunately, the large number of transiting planets  and eclipsing binaries in the NASA \\kep\\ mission target field \\citep{prsa11},  together with the exquisite photometric precision produced by the \\kep\\ space telescope, will provide many opportunities for these tests in the near future."
    },
    "1208/1208.3282_arXiv.txt": {
        "abstract": "{Several clusters of red supergiants have been discovered in a small region of the Milky Way close to the base of the Scutum-Crux Arm and the tip of the Long Bar. Population synthesis models indicate that they must be very massive to harbour so many supergiants. Amongst these clusters, Stephenson~2, with a core grouping of 26 red supergiants, is a strong candidate to be the most massive cluster in the Galaxy.} {Stephenson~2 is located close to a region where a strong over-density of red supergiants had been found. We explore the actual cluster size and its possible connection to this over-density.} {Taking advantage of Virtual Observatory tools, we have performed a cross-match between the DENIS, USNO-B1 and 2MASS catalogues to identify candidate obscured luminous red stars around Stephenson~2, and in a control nearby region. More than 600 infrared bright stars fulfill the colour criteria, with the vast majority having a counterpart in the $I$ band and $>400$ being sufficiently bright in $I$ to allow observation with a 4-m class telescope. We have observed a subsample of $\\sim$250 stars, using the multi-object, wide-field, fibre spectrograph AF2 on the WHT telescope in La Palma, obtaining intermediate-resolution spectroscopy in the 7500--9000\\AA\\ range. We derive spectral types and luminosity classes for all these objects and measure their radial velocities.} {Our targets turn out to be G and K supergiants, late ($\\geq$M4) M giants, and M-type bright giants (luminosity class II) and supergiants. We find $\\sim$35 red supergiants with radial velocities similar to Stephenson~2 members, spread over the two areas surveyed. In addition, we find $\\sim$40 red supergiants with radial velocities incompatible in principle with a physical association.} {Our results show that Stephenson~2 is not an isolated cluster, but part of a huge structure likely containing hundreds of red supergiants, with radial velocities compatible with the terminal velocity at this Galactic longitude (and a distance $\\sim$6~kpc). In addition, we find evidence of several populations of massive stars at different distances along this line of sight.} ",
        "introduction": "In the past few years, a flurry of discoveries has revealed several clusters of red supergiants (RSGs) located in a small region of the Galactic plane, between $l=24\\degr$ and $l=29\\degr$ \\citep[e.g.,][]{figer06, davies07, clark09,neg11}. These clusters are so heavily obscured  that, until now, the only members observed are the RSGs. Population synthesis models suggest that the clusters must contain very large stellar populations to harbour so many RSGs  \\citep[e.g.,][]{davies07}. Recent simulations indicate that, at typical ages $\\tau=10-20\\:$Myr \\citep{davies07,davies08}, one should expect $\\sim$$10^{4}\\,M_{\\sun}$ in stars for each 3--5 RSGs \\citep{simonw51}.  \\defcitealias{davies07}{D07} \\defcitealias{neg11}{Paper~I}  RSGC2 = Stephenson~2 (Ste~2; $l=26\\fdg2$, $b=-0\\fdg1$) is the least obscured and apparently most massive of the red supergiant clusters. Discovered by \\citet{stephenson}, the cluster was found to have a clump of RSGs \\citep{ortolani}. \\citet[henceforth D07]{davies07} obtained $K$-band spectroscopy of a large number of bright sources within $7\\arcmin$ of the nominal cluster centre, finding more than 25 RSGs that shared similar radial velocities. After defining an average radial velocity for the cluster and eliminating outliers, they found $v_{{\\rm LSR}}=+109.3\\pm0.7\\:{\\rm km}\\,{\\rm s}^{-1}$, with the uncertainty representing Poisson statistics for 26 likely members with measured values \\citepalias{davies07}. Stars were considered members if their $v_{{\\rm LSR}}$ was within $\\pm10\\:{\\rm km}\\,{\\rm s}^{-1}$ of the average value. On the other hand, \\citet{deguchi} obtain $v_{{\\rm LSR}}$\\,$=+96.5\\:{\\rm km}\\,{\\rm s}^{-1}$ from measurements of SiO masers associated with four cluster members. They attribute the difference to systematic effects.  The age of Ste~2 has been estimated at  $\\tau$\\,$=17\\pm3$~Myr \\citepalias{davies07}, consistent with a value $\\sim$20~Myr derived by \\citet{ortolani}. Assuming that the 26 RSGs are members, \\citetalias{davies07} estimate, using population synthesis models, that the underlying cluster population must be $M_{{\\rm cl}}\\ga5\\times10^{4}\\,M_{\\sun}$. Conversely, \\citet{deguchi} claim that the spatial distribution of RSGs argues for the existence of two clusters, which they call Ste~2 and Ste~2 SW, though their radial velocities do not seem to be significantly different. In summary, the actual spatial extent, and hence total mass, of Ste~2 is still poorly determined.  The line of sight to the RSG clusters (RSGCs) passes through several Galactic arms and reaches an area of very high obscuration at a distance comparable to those of the clusters \\citep{neg10}. Stellar densities are extremely high at moderate $K$ magnitudes, and membership of individual stars is difficult to assess. \\citetalias{davies07} found a large population of field stars with $K_{{\\rm S}}$ magnitudes comparable to the supergiants. Indeed, down to their limiting magnitude $K_{{\\rm S}}=8$, there are many more field stars than RSG members in the $r=7\\arcmin$ circle analysed (though RSGs are more numerous than field stars for $K_{{\\rm S}}<6.2$).  \\citetalias{davies07} identify these field stars as predominantly foreground giants, though several putative RSGs with $v_{{\\rm rad}}$ suggestive of non-membership were found.  In a recent paper (\\citealt{neg11}; from now on, \\citetalias{neg11}), we studied the spatial extent of another cluster of red supergiants, RSGC3, with estimated $\\tau$=16--20$\\,$Myr and an inferred $M_{{\\rm cl}}$=2--4$\\times10^4\\,M_{\\sun}$  \\citep{clark09,alexander09}, finding that it is part of an association, which includes smaller clusters of red supergiants. This association contains $\\ga 50$ RSGs and hence a total mass $\\ga 10^{5}\\,M_{\\sun}$. This size has been confirmed (and likely surpassed) by the recent discovery of another cluster of RSGs in its vicinity \\citep{carlos}.  The connection of the RSGC3 association to Ste~2 and the other RSG clusters is unclear. Though all the clusters are located at a similar distance and have similar ages, they span $\\sim$500~pc in projection. There have been suggestions that this Scutum Complex represents a giant star formation region triggered by the dynamical excitation of the Galactic Bar, whose tip is believed to intersect the Scutum-Crux Arm close to this region (see discussion in \\citetalias{davies07}; \\citealt{garzon}). The possibility of a hierarchical star-forming structure extending over such a long distance seems unlikely as it would be larger than the largest cluster complexes seen in external galaxies \\citep[for instance, $\\sim240$~pc in M51;][]{bastian}, even though less coherent structures \\citep[aggregates in the terminology of][]{efremov04} are seen spanning comparable distances.  All the RSG clusters have similar ages, in the 10--20~Myr range, but this does not guarantee an association. RSGC1, with an age $\\sim$12~Myr \\citep{davies08}, could be more closely allied to its close neighbour, the Quartet cluster, with an age between 3--8~Myr \\citep{messi09}, while the RSGC3 association lies very close in the sky to the W43 star-forming complex, also at $d$\\,$\\sim6$~kpc, which covers the $l=29\\fdg5-31\\fdg5$ range \\citep{nguyen11}.  The distribution of Galactic \\ion{H}{ii} regions shows a marked maximum towards $l\\sim30\\degr$--$31\\degr$ \\citep{bania, anderson11}, where tens of sources display radial velocities in the $v_{{\\rm LSR}}\\approx90-105\\:{\\rm km}\\,{\\rm s}^{-1}$ interval. This is generally interpreted as the tangent point of the Scutum-Crux arm. There is a secondary maximum towards $l\\sim24\\degr$, where many sources have $v_{{\\rm LSR}}\\approx110\\:{\\rm km}\\,{\\rm s}^{-1}$, similar to the systemic velocity of RSGC1 \\citep[$l$\\,$=25\\fdg3$; $v_{{\\rm LSR}}=123\\:{\\rm km}\\,{\\rm s}^{-1}$;][]{davies08}. There is a local minimum towards $l\\sim26\\degr$, though the number of \\ion{H}{ii} regions in that direction is still high \\citep{anderson11}. Galactic molecular clouds seem to follow a similar distribution \\citep[e.g.,][]{rathborne09}.  In this paper, we study the distribution of luminous red stars in the vicinity of Ste~2. Though the cluster core is clearly defined by a strong concentration of bright infrared sources, its actual extent is poorly determined. The spatial distribution of known members is very elongated, and the possibility that the cluster is surrounded by a halo has not been studied. Likewise, the fact that RSGC3 is surrounded by an extended association opens the possibility that Ste~2 is not isolated either.  Following the strategy outlined in \\citetalias{neg11}, we have used a multi-object spectrograph with high multiplexing capability to obtain intermediate-resolution spectra of a large sample of bright infrared sources in the spectral region around the \\ion{Ca}{ii} triplet. In Section~\\ref{sec:target}, we discuss the photometric criteria used for selection of candidates and describe the observations of a sample of these candidates in the region around Ste~2 and a second sample in a control area about $1\\degr$ away. In Section~\\ref{sec:calib}, we summarise the analysis methodology. We have then used the criteria and calibrations developed in \\citetalias{neg11} to assign spectral types and calculate colour excesses. We also measure radial velocities for all our sources. The results of the analysis are presented in Section~\\ref{sec:res}. Finally, we discuss their implications in Section~\\ref{sec:discu} and wrap up with the conclusions. ",
        "conclusions": "In this paper, we have presented the results of a spectroscopic survey of red luminous stars along the Galactic Plane in the $l=26\\degr$--$28\\degr$ region, covering the vicinity of the starburst cluster Stephenson~2 and a control region $1\\fdg1$ away. We have observed $\\sim$250~stars amongst $>600$ candidates ($>400$ actually observable with our instrumentation), finding that all except one are red luminous stars. We find $\\sim$80 red supergiants, most of M spectral type. This number represents a significant fraction of the RSGs known in the Galaxy, in spite of the fact that we have only surveyed $\\sim$1.7~deg$^{2}$, and checked $\\sim$40\\% of the candidates. The capability to select red luminous stars with very simple colour cuts has been fully demonstrated.  The RSGs found are not strongly concentrated towards Ste~2 (we stress that our observation technique does not allow us to sample but a very small fraction of densely concentrated populations), but seem to be rather uniformly spread over the area surveyed. Moreover, about half the RSGs have radial velocities and reddenings that make a physical connection with the cluster very unlikely.  On the one hand, these findings demonstrate beyond any reasonable doubt that the clusters of RSGs are not isolated, but rather part of extended stellar associations, a result that had already been found for RSGC3 (\\citetalias{neg11}; \\citealt{carlos}). This diffuse population likely extends over several squared degrees and comprises several hundred objects. Based on their reddenings and radial velocities, members of this population cannot be located at distances from the Earth significantly different from those of the RSGCs. Perhaps they represent the result of starburst caused by the tip of the Galactic bar, as suggested by \\citet{lopez}. Wider-field searches can confirm the existence of such a large population and probe its spatial extent.  On the other hand, the high number of RSGs not associated with Ste~2 bears witness to the richness of the young stellar population towards the inner Galaxy. Even though the tip of the Galactic bar may create special conditions at the distance of the Scutum Complex, the foreground population along this sightline is unlikely to be affected. If so, such high numbers could be representative of {\\em all} sightlines towards the inner Milky Way.  Our results show that the limits of Stephenson~2 cannot be clearly defined. A compact core, with radius $\\la2\\farcm5$ and containing about 20 RSGs, is surrounded by an extended association that merges into a general background with an important over-density of RSGs. To investigate the spatial extent of this overdensity and test whether it covers the whole region containing the clusters of red supergiants ($l\\sim 24\\degr - 29\\degr$), a wider-field spectroscopic survey is needed. Such a survey is currently underway and will be presented in a future paper.  Finally, the technique used here offers a novel way of tracing Galactic structure. Red supergiants are very bright objects in the infrared and can be seen through heavy extinction. Even though some of them may have peculiar velocities, their large numbers and their membership in clusters make them very good tracers. Small clusters containing $\\sim$5 RSGs must be relatively common in areas of high young population, and they can serve as very good tracers of Galactic structure. Even if they are too reddened to be accessible around the \\ion{Ca}{ii} triplet, they are very good targets for moderately-high-resolution observations in the $K$ band. Investigation of further photometric criteria, making use of a combination of 2MASS and GLIMPSE colours, seems the next natural step along this approach to penetrating into the inner Milky Way."
    },
    "1208/1208.6234_arXiv.txt": {
        "abstract": "Several independent measurements have confirmed the existence of fluctuations ($\\delta F_{\\rm obs}\\approx 0.1~\\rm nW/m^{2}/sr$ at $3.6~\\rm \\mu m$) up to degree angular scales in the source-subtracted Near InfraRed Background (NIRB) whose origin is unknown. By combining high resolution cosmological N-body/hydrodynamical simulations with an analytical model, and by matching galaxy Luminosity Functions (LFs) and the constraints on reionization simultaneously, we predict the NIRB absolute flux and fluctuation amplitude produced by high-$z$ ($z > 5$) galaxies (some of which harboring Pop III stars, shown to provide a negligible contribution). This strategy also allows us to make an empirical determination of the evolution of ionizing photon escape fraction: we find $f_{\\rm esc} = 1$ at $z \\ge 11$, decreasing to $\\approx 0.05$ at $z = 5$.  In the wavelength range $1.0-4.5~\\rm \\mu m$, the predicted cumulative flux is  $F =0.2-0.04~\\rm nW/m^2/sr$. However, we find that the radiation from high-$z$ galaxies (including those undetected by current surveys) is insufficient to explain the amplitude of the observed fluctuations: at $l=2000$, the fluctuation level due to $z > 5$ galaxies is $\\delta F = 0.01-0.002~\\rm nW/m^2/sr$, with a relative wavelength-independent amplitude $\\delta F/F = 4$\\%. The source of the missing power remains unknown. This might indicate that an unknown component/foreground, with a clustering signal very similar to that of high-$z$ galaxies, dominates the source-subtracted NIRB fluctuation signal. ",
        "introduction": "\\label{methods} \\subsection{The Absolute Flux}  At $z=0$, the cumulative flux of the NIRB observed at frequency $\\nu_0$ is the integrated contribution of sources whose emission is shifted into a band of central frequency $\\nu_0$. Following \\cite{2006MNRAS.368L...6S}, we write it as \\begin{eqnarray} F&=&\\nu_0I_{\\nu_0} \\nonumber \\\\ &=&\\nu_0\\int_{z_{\\rm min}}^{z_{\\rm max}}\\epsilon(\\nu,z){\\rm e}^{-\\tau_{\\rm eff}(\\nu_0,z)}\\frac{dr_p}{dz}dz \\nonumber\\\\ &=&\\int_{z_{\\rm min}}^{z_{\\rm max}}cdz\\frac{\\nu\\epsilon(\\nu,z){\\rm e}^{-\\tau_{\\rm eff}(\\nu_0,z)}}{H(z)(1+       z)^2}, \\label{I0} \\end{eqnarray} where $r_p$ is the proper distance, $\\nu=(1+z)\\nu_0$ is the rest frame frequency, $\\epsilon(\\nu,z)$ is the comoving specific emissivity, $H(z)$ is the Hubble parameter given by $H(z)=H_0\\sqrt{\\Omega_m(1+z)^3+\\Omega_\\Lambda}$ in a flat $\\Lambda$CDM cosmology, $c$ is the speed of light. The effective optical depth of absorbers between redshift 0 and $z$, $\\tau_{\\rm eff}$, is composed of two parts: the line absorption and the continuum absorption; we use the expressions in \\cite{2003MNRAS.339..973S}.  We calculate the emissivity (see also Appendix for further discussions on subtleties related to the various approximations used in the literature) from the results of the simulation presented in \\cite{2011MNRAS.414..847S}, which includes a detailed treatment of chemical enrichment developed by \\cite{2007MNRAS.382.1050T}. In our model, both Pop II stars and Pop III stars are assumed to follow the Salpeter initial mass function (IMF) \\citep{1955ApJ...121..161S}, for Pop II stars the mass range is $0.1-100~M_\\odot$, while for Pop III stars the mass range is set to be $100-500~M_\\odot$. Some recent works indicate that Pop III stars may not be so massive as was predicted previously, but may be limited to $\\lesssim50~M_\\odot$ \\citep{2011Sci...334.1250H}. Our choice then corresponds to an upper limit to the contribution of these sources. Using this simulation, \\cite{2011MNRAS.414..847S} generated the LFs of galaxies down to the magnitude far below the current observation limits at high redshifts. In the redshift range $5 < z < 10$, the simulated LFs match the observed ones almost perfectly in the overlapping luminosity range.  Suppose the specific luminosity of the $i$-th galaxy in the simulation box is $L^i_\\nu(z)$ at redshift $z$, the comoving specific emissivity is then \\footnote{The dust absorption in the host galaxy may lower its luminosity in the rest frame UV band, therefore reduce their contribution to the NIRB. However, as shown by \\cite{2011MNRAS.414..847S}, for high redshift dwarf galaxies considered by us here, this effect is negligible. The absorption is possible important for very massive objects in the bright end of LFs, but their contribution to the emissivity is very small. The lack of significant dust absorption is further supported by the observed very blue UV-continuum slopes of high-$z$ galaxies reported in \\cite{2009ApJ...705..936B}. Therefore, we ignore the effect of dust in current calculations. Furthermore, we will show latter that the high-$z$ galaxies are inadequate to explain the amplitude of the observed fluctuations, considering the dust absorption would only strength this point. } \\begin{equation} \\epsilon(\\nu,z)=\\frac{1}{4\\pi}\\frac{\\sum_{i=1}^NL^i_\\nu(z)}{V}, \\label{e} \\end{equation} where $V$ is the comoving volume of the simulation, $N$ is the total number of galaxies in the simulation box at redshift $z$.  In the emissivity calculation, we must correct for rare bright galaxies that are not caught by the simulation due to the finite box size ($10~h^{-1}$Mpc). We follow the steps in \\cite{2011MNRAS.414..847S}. We first calculate the absolute magnitude corresponding to the mean luminosity of the two brightest galaxies in the simulation box, $M_{\\rm UV,up}$. The contribution (to be added to the numerator in Eq. (\\ref{e})) from galaxies brighter than this magnitude is obtained by integration \\begin{equation} L^{\\rm corr}_\\nu(z)=V\\int_{-25}^{M_{\\rm UV,up}} L^1_\\nu(z)\\frac{L_{\\rm UV}}{L^1_{\\nu_{\\rm UV}}(z)}\\Phi(M_{\\rm UV},z) dM_{\\rm UV}, \\label{corr} \\end{equation} where $L^1_\\nu(z)$ is the luminosity of the brightest galaxy in the simulation (we assume all rare bright galaxies have the same Spectral Energy Distribution (SED) of this one), $L_{\\rm UV}$ is the luminosity corresponding to the UV absolute magnitude $M_{\\rm UV}$. The wavelength used to calculate the absolute UV magnitude in this paper is $1700~\\rm \\AA$, $\\nu_{\\rm UV}$ is the frequency corresponding to this wavelength. In observations, the selected wavelength corresponding to the UV absolute magnitude may be somewhat different in different measurements and at different redshifts \\citep{2007ApJ...670..928B,2010ApJ...709L..16O,2010ApJ...709L.133B}, however, our results are not sensitive to such differences. For the LF $\\Phi(M_{\\rm UV},z)$ in the redshift range $5 < z < 10$, we use the Schechter formula \\citep{1976ApJ...203..297S} with the redshift-dependent parameters given by \\cite{2011ApJ...737...90B} (see their Sec. 7.5), who fitted the observed LFs in $z\\sim4-8$ and extrapolated them to higher redshifts. For redshifts above 10, we simply add an exponential tail normalized to  the simulated LF amplitude at $M_{\\rm UV,up}$. We find that this results only a small correction. As discussed below (see also bottom panel of Figure \\ref{Inu0}), $\\sim 90\\%$ of the high-$z$ galaxy contribution to the NIRB flux comes from sources at $5<z<8$  where the correction is at most $12\\%$. This correction is also applied to the calculation of ionizing photons below.  For each galaxy, the radiation comes from two different mechanisms: the stellar emission and the nebular emission. The former comes directly from the surface of stars, while the latter is generated by the ionized nebula around stars and depends on the fraction of ionizing photons that cannot escape into the intergalactic medium (IGM), i.e., $1-f_{\\rm esc}$, where $f_{\\rm esc}$ is the escape fraction. Ionizing photons escaping from galaxies would ionize the IGM; such ionized gas could also produce the nebular emission. However, due to the very low recombination rate, as shown in, e.g., \\cite{2001MNRAS.321..593N} and \\cite{2012ApJ...756...92C}, its emissivity is much weaker than the radiation from galaxies, so we ignore this contribution in this paper. The IGM contribution to the NIRB fluctuations is also negligible \\citep{2010ApJ...710.1089F}.   To determine the escape fraction averaged over the galaxy populations present at a given redshift, we proceed as follows. First we compute the number of ionizing photons emitted per baryon in collapsed objects as: \\begin{equation} f_\\star N_\\gamma= \\frac{\\sum_{i=1}^N \\left[q^{\\rm II}_{\\rm H}(\\tau^{{\\rm II},i},Z^i)\\dot{M}^{{\\rm II},i}_\\star\\tau^{{\\rm II} ,i}+q^{\\rm III}_{\\rm H}M^{{\\rm III},i}_\\star\\tau^{\\rm III}\\right]}{\\sum_{i=1}^NM^i_{\\rm gas}}, \\end{equation} where  $q^{\\rm II}_{\\rm H}$ is the emission rate of ionizing photos from Pop II stars  (this quantity depends on both the age and the metallicity of the stellar population)  corresponding to a continuous star formation rate $1~M_\\odot\\rm yr^{-1}$. We derive this quantity from the {\\tt Starburst99} templates\\footnote{http://www.stsci.edu/science/starburst99/docs/default.htm} \\citep{1999ApJS..123....3L,2005ApJ...621..695V,2010ApJS..189..309L} adopting the mean age, $\\tau^{{\\rm II},i}$, and metallicity, $Z^i$, of each simulated galaxy.  $q^{\\rm III}_{\\rm H}$ is the emission rate of ionizing photons for Pop III stars according to \\cite{2002A&A...382...28S} \\footnote{Note the different dimensions of $q^{\\rm II}_{\\rm H}$ and $q^{\\rm III}_{\\rm H}$: the former corresponds to per unit star formation rate while the latter corresponds to per unit stellar mass.}. $M^i_{\\rm gas}$ is the gas content, $\\dot{M}^{{\\rm II},i}_\\star$ is the mean star formation rate of Pop II stars in this galaxy, while $M^{{\\rm III},i}_\\star$ is the cumulative mass of Pop III stars. We use a mean lifetime $\\tau^{\\rm III}=2.5\\times10^6$~yr for massive Pop III stars \\citep{2002A&A...382...28S,2011MNRAS.414..847S}.  We then compare the above quantity with the number of ionizing photons per baryon in collapsed objects, $N_{\\rm ion}$, required by interpreting the observations as in \\cite{2012MNRAS.419.1480M} (the ``mean\" value) to get the escape fraction, i.e., $f_{\\rm esc}={\\rm min}(\\frac{CN_{\\rm ion}}{f_\\star N_\\gamma},1.0)$, where $C$ is the clumping factor. Throughout this paper we assume $C = 1$ to get the minimum $f_{\\rm esc}$ therefore the maximum contribution of the nebular emission to the NIRB. We note that the clumping factor could be higher than 1 even at high redshifts \\citep{2009MNRAS.394.1812P,2012ApJ...747..100S}. For example, \\cite{2012ApJ...747..100S} gives $C\\approx 3~(1.7)$ at $z = 5~(9)$ by numerical simulations. Our nebular emission is therefore reduced by about 10\\% to 60\\% from $z = 5$ to $z = 9$ if this clumping factor is adopted. However, for Pop II stars which are the dominant contributors to the NIRB, the nebular emission is much smaller than the stellar emission. So the final reduction in the NIRB would be much smaller. Furthermore, we will show later that the flux from high redshift galaxies and Pop III stars is unable to explain the observed fluctuations level, so that a reduction in the NIRB flux would in any case strengthen this conclusion. We plot the derived $f_{\\rm esc}$ as a function of redshift in Figure \\ref{fesc}. There is a clear trend of an increasing escape fraction towards higher redshifts; it reaches 1 at $z \\approx 11$. At $z = 5$, the final redshift of the simulation, $f_{\\rm esc}\\approx0.05$. Although required by reionization data, an increasing trend of $f_{\\rm esc}(z)$ has not yet been fully understood theoretically in spite of the several, often conflicting, studies on this problem.   Based on observations, \\cite{2006MNRAS.371L...1I} concluded that $f_{\\rm esc} > 0.1$ when $z > 4$; by combining the observations of Lyman $\\alpha$ absorption and UV LF, and also using N-body simulations and semi-analytical prescriptions to model the ionizing background, \\cite{2010PASA...27..110S} found that for galaxies at $z\\sim5.5-6$, if the minimum mass of star forming galaxies corresponds to the hydrogen cooling threshold, $f_{\\rm esc}\\sim0.05-0.1$; \\cite{2010MNRAS.401.2561W} used the star formation rate derived from gamma-ray burst observations to conclude that in the redshift range $4-8.5$, $f_{\\rm esc}\\sim0.05$; \\cite{2009ApJ...693..984W}, by radiation hydrodynamical simulations, found that at redshift 8 for galaxies with $M_{\\rm vir} < 10^{7.5}~M_\\odot$, $f_{\\rm esc}\\sim0.05-0.1$, while for more massive galaxies $f_{\\rm esc}\\sim0.4$, if a normal IMF is adopted; also via simulations, \\cite{2010ApJ...710.1239R} found $f_{\\rm esc}\\sim0.8$ when $z=10$. The escape fraction derived by us is broadly consistent with these values. The important difference however is that our derivation of $f_{\\rm esc}$ matches both the LF and the reionization history simultaneously, i.e., a more phenomenological derivation, so that we can get around the detailed physical mechanisms of the escape fraction. Different from our approach, \\cite{2013MNRAS.428L...1M} computed the LFs of high redshift galaxies by means of semi-analytical models and derived the star formation efficiency $f_\\star$ required to match the observed ones. They found $f_{\\rm esc} \\approx 0.07$ at $z = 6$ and $f_{\\rm esc} \\approx 0.16$ at $z = 7$, which are consistent with our $f_{\\rm esc} \\approx 0.06~(0.18)$ at those two redshifts. At higher redshifts, however, their escape fraction is somewhat lower than ours.  As a final remark, we underline that when computing the escape fraction, we do not make a distinction between Pop III and Pop II stars. In the calculation of $f_\\star N_{\\gamma}$ ionizing photons from both populations are accounted for, $f_{\\rm esc}$ can be regarded as a kind of ``effective\" escape fraction averaged over the galaxy population. In principle, $f_{\\rm esc}$  for Pop III stars should be higher due to their harder spectrum. However, as we will see in Section \\ref{results}, Pop III stars only contribute a negligible flux to the present-day NIRB, a more detailed modeling is then not necessary.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{./fig1.eps} \\caption{Escape fraction evolution from joint LF-reionization constraints.} \\label{fesc} \\end{center} \\end{figure}    The luminosity of the $i$-th galaxy, $L^i_\\nu$, is the sum of the contribution of Pop II and Pop III stars. For Pop II stars we use the age and metallicity dependent spectrum templates provided by the {\\tt Starburst99} code. The nebular emission contribution has been renormalized by adopting the escape fraction computed above. In addition to the free-free, free-bound and two-photon emissions which have already been included in {\\tt Starburst99}, we add the Lyman $\\alpha$ emission to the template by using \\citep{2006ApJ...646..703F} \\begin{equation} l_{\\alpha}(\\nu,\\tau^{{\\rm II},i},Z^i,z)=f_{\\alpha}h_{\\rm p}\\nu_{\\alpha}\\phi(\\nu-\\nu_\\alpha)q^{{\\rm II}}_{\\rm H}( \\tau^{{\\rm II},i},Z^i)[1-f_{\\rm esc}(z)], \\label{lya} \\end{equation} in which $f_\\alpha=0.64$ \\citep{2006ApJ...646..703F}, $h_{\\rm p}$ is the Plank constant, $\\nu_\\alpha=2.47\\times10^{15}$~Hz is the frequency of Lyman $\\alpha$ photons. We use the line profile $\\phi(\\nu-\\nu_\\alpha)$ provided in \\cite{2002MNRAS.336.1082S}: \\begin {equation} \\phi(\\nu-\\nu_\\alpha)= \\left\\{ \\begin{tabular}{ll} $\\nu_\\star(z)(\\nu-\\nu_\\alpha)^2{\\rm exp}[-\\nu_\\star(z)/|\\nu-\\nu_\\alpha|]$~if $\\nu \\le \\nu_\\alpha$ \\\\ 0~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~if $\\nu > \\nu_\\alpha$, \\end{tabular} \\right. \\end{equation} where \\begin{equation} \\nu_\\star(z)=1.5\\times10^{11}\\left(\\frac{\\Omega_bh^2}{0.019}\\right)\\left(\\frac{h}{0.7}\\right)^{-1}\\frac{(1+z)^3}{\\sqrt{\\Omega_m(1+z)^3+\\Omega_\\Lambda}}~\\rm Hz \\end{equation} is the fitted form of results given in \\cite{1999ApJ...524..527L}.   For the template of Pop III stars, $l_\\nu^{\\rm III}$, we still use the spectrum in \\cite{2002A&A...382...28S}, but renormalize the nebular emission part by the factor $1-f_{\\rm esc}$. The luminosity of the $i$-th galaxy is then given by \\citep{2011MNRAS.414..847S} \\begin{equation} L^i_\\nu(z)=l_\\nu^{{\\rm II}}(\\tau^{{\\rm II},i},Z^i,z)\\dot{M}^{{\\rm II},i}_\\star+ l_\\nu^{{\\rm III}}(z)\\dot{M}_\\star^{{\\rm III},i}\\tau^{{\\rm III}}, \\end{equation} here the Lyman $\\alpha$ emission in Eq. (\\ref{lya}) has already been included in $l_\\nu^{{\\rm II}}$.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{./fig2.eps} \\caption{The emissivity of simulated galaxies at redshift 12.0 (dash-dotted), 9.0 (dashed) and 6.0 (solid) respectively.} \\label{emiss} \\end{center} \\end{figure}  With the luminosity for each galaxy given as above, we can then obtain the emissivity according to Eq. (\\ref{e}). As an example, we plot the $\\nu\\epsilon(\\nu,z)$ at redshifts 12.0, 9.0 and 6.0 respectively in Figure \\ref{emiss}. At high redshifts, the escape fraction $\\approx1.0$, yielding a very weak Ly$\\alpha$ line, since such emission is produced by recombinations of the ionized nebula around stars. At lower redshifts, the escape fraction drops, while more ionizing photons are absorbed by the material around the stars, and producing more Ly$\\alpha$ emission which is more clearly seen in the spectrum.  The part of spectrum with energy below 10.2 eV is of the most interest to us, here the spectrum becomes increasingly flatter at later time. For example, at $z = 12$, the slope of $\\nu\\epsilon(\\nu,z) \\propto \\nu^\\beta$ with $\\beta \\approx 2$, while at $z = 5$ $\\beta \\approx 1.2$. This is clearly the result of an aging effect enhancing the rest frame optical/IR bands flux with respect to the UV ones. Since the NIRB from $z > 5$ galaxies is dominated by the lower redshift galaxies ($5 < z < 8$), we do not expect to have a very steep NIRB spectrum, as we will see in the results presented in Sec. \\ref{results}.    \\subsection{NIRB Fluctuations} Using the Limber approximation, the angular power spectrum of the fluctuations of the flux field is \\citep{2012ApJ...756...92C} \\begin{equation} C_l=\\int_{z_{\\rm min}}^{z_{\\rm max}}\\frac{dz}{r^2(z)(1+z)^4}\\frac{dr}{dz}[\\nu\\epsilon(\\nu,z){\\rm e}^{-\\tau_{\\rm eff}(\\nu_0,z)}]^2P_{\\rm gg}(k,z), \\label{angpow} \\end{equation} where $r(z)$ is the comoving distance and $P_{\\rm gg}(k,z)$ is the galaxy-galaxy power spectrum, $k=l/r(z)$. In Eq. (\\ref{angpow}) we assume that the luminous properties of galaxies are independent of their locations, so that the only factor which determines their contribution to the NIRB fluctuations is their spatial fluctuations (see \\citealt{2012MNRAS.421.2832S} for an improved model).  The $10~h^{-1}$~Mpc box size of \\cite{2011MNRAS.414..847S} simulations is too small to provide us with the large scale correlation function of galaxies (for sources at $z = 6$, the comoving transverse separation corresponding to $1^{\\ensuremath{\\circ}}$ angular size is about $100~h^{-1}$Mpc), so we use the halo model \\citep{2002PhR...372....1C,2004MNRAS.348..250C} to calculate the galaxy-galaxy power spectrum. This power spectrum is composed of two parts, the one-halo term from the correlation of galaxies in the same halo (including the central galaxies and satellite galaxies), and the two-halo term from galaxies in different halos: \\begin{equation} P_{\\rm gg}(k,z)=P^{1h}_{\\rm gg}(k,z)+P^{2h}_{\\rm gg}(k,z). \\end{equation} Assuming that the distribution of galaxies in a halo traces the profile of dark matter, and the mean number of central galaxies and satellite galaxies in a halo with mass $M$ are $\\langle N_{\\rm sat}\\rangle$ and $\\langle N_{\\rm cen}\\rangle$, respectively, we have \\begin{equation} P^{1h}_{\\rm gg}(k,z)= \\int_{M_{\\rm min}(z)}^{M_{\\rm max}(z)}dM\\frac{dn}{dM} \\frac{2\\langle N_{\\rm sat}\\rangle \\langle N_{\\rm cen}\\rangle u(M,k)+ \\langle N_{\\rm cen}\\rangle^2 u^2(M,k)}{\\bar{n}^2_{\\rm gal}}, \\end{equation} and \\begin{equation} P^{2h}_{\\rm gg}(k,z)=P_{\\rm lin}(k,z) \\times\\left[\\int_{M_{\\rm min}(z)}^{M_{\\rm max}(z)} dM\\frac{dn}{dM}b(M,z) \\frac{\\langle N_{\\rm sat}\\rangle+\\langle N_{\\rm cen} \\rangle}{\\bar{n}_{\\rm gal}}u(M,k)\\right]^2. \\end{equation} In the above expressions, $M_{\\rm min}(z)$ is the minimum mass of halos that could host galaxies, and we set it to be the minimum mass of halos that contain stars in our simulations, which is $\\sim(2-8)\\times10^7~M_\\odot$, depending on the redshift. $M_{\\rm max}(z)$ is the maximum mass contributing to the emissivity and the clustering, we will describe how to determine it later, ${dn}/{dM}$ is the mass function \\citep{1999MNRAS.308..119S,2001MNRAS.323....1S}, while $u(M,k)$ is the normalized Fourier transform of the halo profile. For a NFW profile \\citep{1997ApJ...490..493N}, the analytical expression is given by \\cite{2002PhR...372....1C}, and we use the concentration parameter, $c_{\\rm M}$, from \\cite{2012MNRAS.423.3018P} which fits simulations well. However, we find that our results are insensitive to the use of different concentrations, as e.g., from \\cite{2011ApJ...736...59Z}. Even if a very different concentration parameter is adopted, its impacts are non-negligible only in the one-halo term which dominates the signal at small scales, where galaxy clustering is well below the shot noise. As we are interested primarily in the large-scale ($> 1'$) clustering, our conclusions are unaffected by the adopted value of $c_{\\rm M}$. Finally, for the halo bias $b(M,z)$ we use the formula and fitted parameters given by \\cite{2010ApJ...724..878T}, which is higher than \\cite{2001MNRAS.323....1S} for massive halos, but is better fit to simulations. The linear matter power spectrum $P_{\\rm lin}(k,z)$ is taken from \\cite{1998ApJ...496..605E}.  The mean number of central galaxies and satellites in a halo with mass $M$ is modeled by the halo occupation distribution (HOD) model \\citep{2005ApJ...633..791Z}, \\begin{equation} \\langle N_{\\rm cen} \\rangle =\\frac{1}{2}\\left[1+{\\rm erf} \\left(\\frac{{\\rm log_{10}}M-{\\rm log_{10}}M_{\\rm min} }{\\sigma_{{\\rm log_{10}}M}}\\right)\\right], \\end{equation} and \\begin{equation} \\langle N_{\\rm sat} \\rangle=\\frac{1}{2}\\left[1+{\\rm erf} \\left(\\frac{{\\rm log_{10}}M-{\\rm log_{10}}2M_{\\rm min} }{\\sigma_{{\\rm log_{10}}M}}\\right)\\right] \\left(\\frac{M}{M_{\\rm sat}}\\right)^{\\alpha_s}. \\end{equation} We adopt the parameters $M_{\\rm sat}=15M_{\\rm min}$, $\\sigma_{{\\rm log_{10}}M}=0.2$ and $\\alpha_s=1.0$, which are from both simulations and semi-analytical models \\citep{2005ApJ...633..791Z}, and observations \\citep{2011ApJ...736...59Z}. With the mean number of central and satellite galaxies in each halo, the galaxy number density is simply \\begin{equation} \\bar{n}_{\\rm gal}=\\int_{M_{\\rm min}(z)}^{M_{\\rm max}(z)} dM\\frac{dn}{dM}\\left( \\langle N_{\\rm cen}\\rangle+\\langle N_{\\rm sat} \\rangle \\right). \\end{equation}   In addition to the above galaxy clustering, Poisson fluctuations in the number of galaxies would generate shot noise in observations, whose power spectrum dominates at small scales. If the redshift derivative of the number of sources with flux between $S$ and $S+dS$ is $\\frac{d^2N}{dSdz}$, the angular power spectrum of such shot noise is \\begin{equation} C^{\\rm SN}_l=\\frac{1}{\\Delta \\Omega}\\int dz \\int dSS^2\\frac{d^2N}{dSdz}= \\frac{1}{\\Delta \\Omega}\\int dz \\int dMS^2\\frac{d^2N}{dMdz}, \\end{equation} where $\\Delta \\Omega$ is the beam angle. Considering that \\begin{equation} \\frac{d^2N}{dMdz}=\\frac{dn}{dM}\\Delta \\Omega r^2\\frac{dr}{dz}, \\end{equation} and \\begin{equation} S=\\frac{L_{\\nu}(M){\\rm e}^{-\\tau_{\\rm eff}(\\nu_0,z)}}{4\\pi r^2(1+z)}, \\end{equation} where $L_{\\nu}(M)$ is the luminosity of halos with mass $M$, the shot noise power spectrum is \\begin{equation} C^{\\rm SN}_l=\\int dz\\frac{{\\rm e}^{-\\tau_{\\rm eff}(\\nu_0,z)}}{r^2(1+z)^2}\\frac{dr}{dz} \\int \\left[\\frac{L_{\\nu}(M)}{4\\pi M}\\right]^2 M^2\\frac{dn}{dM}dM. \\end{equation} As we assume the luminous properties of galaxies are independent of their location, in the square bracket we can simply use an average light-to-mass ratio that is independent of the halo mass, \\begin{equation} \\frac{1}{4\\pi}\\left.\\int L_{\\nu}(M)\\frac{dn}{dM}dM\\middle/ \\int M\\frac{dn}{dM}dM\\right.=\\frac{\\epsilon(\\nu,z)}{\\rho_h}. \\end{equation} We finally obtain the shot noise angular power spectrum \\begin{equation} C^{SN}_l=\\int_{z_{\\rm min}}^{z_{\\rm max}}\\frac{cdz}{H(z)r^2(z)(1+z)^4}P^{\\rm SN}(z), \\end{equation} where \\begin{equation} P^{SN}(z)=\\left[\\frac{\\nu\\epsilon(\\nu,z){\\rm e}^{-\\tau_{eff}}} {\\rho_{h}}\\right]^2\\int_{M_{\\rm min}(z)}^{M_{\\rm max}(z)} dMM^2\\frac{dn}{dM}, \\end{equation} and the halo mass density is $$ \\rho_h=\\int_{M_{\\rm min}(z)}^{M_{\\rm max}(z)} dM M\\frac{dn}{dM}$$.  In observations, the detected sources are generally removed down to a certain limiting magnitude, $m_{\\rm lim}$, the residue is the source-subtracted NIRB fluctuations. To simulate this, we also remove bright galaxies in the simulation box and in the bright-end. In theoretical calculations, this limiting magnitude is determined by letting the predicted shot noise level match the values found in the measurements.   The apparent limiting magnitude (at wavelength $\\lambda_0$) is converted into the rest frame absolute magnitude (at wavelength $\\lambda_0/(1+z)$) $M_{\\lambda_0/(1+z)}$ by \\begin{equation} M_{\\lambda_0/(1+z)}=m_{\\rm lim}-DM(z)+2.5 {\\rm log_{10}}(1+z), \\end{equation} where $DM(z)$ is the distance modulus \\citep{2012ApJ...752..113H}. By using a light-to-mass ratio constructed from the simulation, we determine the maximum halo mass $M_{\\rm max}(z)$. Things are slightly more complicated when calculating the absolute flux and the spectrum of fluctuations, since they both depend on wavelength, while the limiting magnitude in observations at different wavelength is different. In this case we simply give the theoretical prediction without removing any sources in the simulation box and the bright-end, as shown in Figure \\ref{Inu0} and Figure \\ref{Cl_spectrum}, corresponding to $M_{\\rm max} = \\infty$. Then we discuss the effects of galaxy removal, i.e., in Figure \\ref{subtract}. Throughout this paper we adopt $z_{\\rm min} = 5$ and $z_{\\rm max} = 19$ unless otherwise specified. ",
        "conclusions": "By combining high resolution cosmological N-body/hydrodynamical simulations and an analytical model, we predicted the contributions to the absolute flux and fluctuations of the NIRB by high redshift ($z > 5$) galaxies, some of which harboring Pop III stars. This is the most robust and detailed theoretical calculation done so far, as we simultaneously match the LFs and reionization constraints. The simulations include the relevant physics of galaxy formation and a novel treatment of chemical feedback, by following the metallicity evolution and implementing the physics of Pop III/Pop II transition based on a critical metallicity criterion. It reproduces the observed UV LFs over the redshift range $5 < z < 10$, and extend it to faint magnitudes far below the detection limit of current observations.  We directly calculate the stellar emissivity from the simulations.  We use {\\tt Starburst99} to generate metallicity and age dependent SED templates, then calculate the luminosity for each galaxy according to its current star formation rate, stellar age and metallicity, instead of using a constant metallicity and average main sequence spectrum template. Except for the mass range of the IMF which has already been fixed in the simulation, there are no other free parameters in the calculation of the emissivity.  By comparing the number of ionizing photons produced per baryon in collapsed objects, $f_\\star N_\\gamma$, in the simulation and the ionizing photon rate $N_{\\rm ion} \\approx f_{\\rm esc} f_\\star N_\\gamma$ deduced from observationally constrained reionization models, we obtained the evolution history of the escape fraction of ionizing photons, $f_{\\rm esc}(z)$. We find $f_{\\rm esc}\\approx 1$ at $z > 11$, decreasing to $\\approx 0.05$ at $z=5$. This escape fraction is used to renormalize the nebular emission of Pop III and Pop II stars in the emissivity.  Pop III stars are unlikely to be responsible for the observed NIRB residual, and their contribution is very small, making up $<1$\\% of the total absolute flux in our calculation. This is the natural result of the much lower star formation rate of Pop III stars compared with Pop II stars in the simulation, since even metals from a single Pop III star could enrich above the critical metallicity a large amount of gas around it \\citep{2007MNRAS.382..945T}. The formation of Pop III stars is regulated by such a chemical feedback mechanism, which limits their contribution to the NIRB. However, a rapid Pop III-Pop II transition brings also a little advantage in terms of integrated emissivity, due to the longer lifetime of Pop II stars \\citep{2012ApJ...756...92C}.   We predict that in the wavelength range $1.0-4.5~\\rm \\mu m$, the NIRB flux from $z > 5$ galaxies (and their Pop III stars) is $\\sim 0.2-0.04~\\rm nW/m^2/sr$, while the fluctuation strength is about $\\delta F = 0.01-0.002~\\rm nW/m^2/sr$ at $l=2000$. If we remove galaxies down to $m_{\\rm lim} = 28$, the above flux level is only slightly reduced; however, by comparing with \\cite{2012ApJ...752..113H}, we find that the flux from $z<5$ dramatically decreases and the remaining becomes comparable to the predicted signal of $z > 5$ galaxies. This implies that in principle it is possible to get the signal from reionization sources by subtracting galaxies down to a certain magnitude.  The relative fluctuation amplitude, $\\delta F/F$, at  $l=2000$ is $\\sim4\\%$, almost independent of the wavelength. This ratio may be helpful to investigate the clustering features of the sources contribute to the NIRB, since the intrinsic properties of galaxies almost cancel out. Despite the difficulties in measuring the absolute flux accurately, it could be treated as a quality indicator in the data reduction process: if a much higher/lower ratio is obtained from the data, this might suggest that a more careful analysis work is required to extract the genuine contribution from reionization sources.  In spite of being accurate and consistent with the observed LFs and reionization data, thus offering a robust prediction of the NIRB contribution from high-$z$ galaxies which likely reionized the universe, a puzzling question remains: the predicted fluctuations are considerably lower than the observed values, indicating that in addition to the contribution from the expected high-$z$ galaxy population (and Pop III stars), we should invoke some other -- yet unknown -- missing component(s) or foreground(s) which dominates the currently observed source-subtracted NIRB. Moreover, the angular clustering of this missing component must be very similar to that of the high redshift galaxies and extends to degree scales. Obviously, this component/foreground must be identified and removed before we are ready to exploit the NIRB to study reionization sources. On the other hand, sources located at $5 < z < 8$ provide about 90\\% of the flux from all sources with $z > 5$ in our simulation; most of them are the faint galaxies currently undetected by deep surveys. Thus, if the above mentioned additional spurious sources/foregrounds can be removed reliably, the NIRB will become the primary tool to investigate the properties of the reionizing sources."
    },
    "1208/1208.1506.txt": {
        "abstract": "{ In this first paper of this series, we present a new approach for studying the chemo-dynamical evolution in disk galaxies, which consists of fusing disk chemical evolution models with compatible numerical simulations of galactic disks. This method avoids known star formation and chemical enrichment problems encountered in simulations. Here we focus on the Milky Way, by using a detailed thin-disk chemical evolution model (matching local observables, which are weakly affected by radial migration) and a simulation in the cosmological context, with dynamical properties close to those of our Galaxy. We examine in detail the interplay between in situ chemical enrichment and radial migration and their impact on key observables in the solar neighborhood, e.g., the age-metallicity-velocity relation, the metallicity distribution, and gradients in the radial and vertical directions. We show that, due to radial migration from mergers at high redshift and the central bar at later times, a sizable fraction of old metal-poor high-[$\\alpha$/Fe] stars can reach the solar vicinity. This naturally accounts for a number of recent observations related to both the thin and thick disks, despite the fact that we use thin-disk chemistry only. Although significant radial mixing is present, the slope in the age-metallicity relation is only weakly affected, with a scatter compatible with recent observational work. While we find a smooth density distribution in the [O/Fe]-[Fe/H] plane, we can recover the observed discontinuity by selecting particles according to kinematic criteria used in high-resolution samples to define the thin and thick disks. We outline a new method for estimating the birth place of the Sun and predict that the most likely radius lies in the range $4.4<r<7.7$~kpc (for a current location at $r=8$~kpc). A new, unifying model for the Milky Way thick disk is offered, where both mergers and radial migration play a role at different stages of the disk evolution. We show that in the absence of early-on massive mergers the vertical velocity dispersion of the oldest stars is underestimated by a factor of $\\sim2$ compared with observations. We can, therefore, argue that the Milky Way thick disk is unlikely to have been formed through a quiescent disk evolution. An observational test involving both chemical and kinematic information must be devised to ascertain this possibility. }  \\titlerunning{Chemodynamical evolution of the Milky Way disk I} \\authorrunning{I. Minchev et al.} ",
        "introduction": "\\label{sec:intro}   Crucial information regarding the dominant mechanisms responsible for the formation of the Milky Way (MW) disk and other Galactic components is encoded in the chemical and kinematic properties of its stars. This conviction has led to unprecedented observational efforts in the past decades, aimed at mapping the chemistry and kinematics of a large number of stars in the MW. Until the end of 2003 most of the information was confined to small local samples for which high-resolution spectroscopic data were obtained (e.g., \\citealt{fuhrmann98} within 25~pc and \\citealt{bensby03} within 100~pc). In 2004, the Geneva Copenhagen Survey \\cite[CGS][]{nordstrom04} obtained the first large spectro-photometric sample of around 16 000 stars as part of a Hipparcos follow-up campaign (hence, also confined to $\\sim$100~pc from the Sun). More recently, optical spectroscopic low-resolution surveys, such as SEGUE \\citep{yanny09} and RAVE \\citep{steinmetz06}, have extended the studied volume to distances of a few kpc from the Sun (with the majority of stars in the distance range 0.5-3~kpc), and increased the numbers of stars with chemo-kinematical information by more than an order of magnitude ($>2\\times10^5$ spectra for SEGUE and $>5\\times10^5$ spectra for RAVE, see \\citealt{steinmetz12}). This effort will be soon complemented by high-resolution spectroscopic surveys both in the optical -- HERMES \\citep{freeman10} and in the near-infrared -- APOGEE \\citep{allende08,majewski10}. APOGEE aims at measuring chemo-kinematic properties of around $10^5$ stars close to the Galactic plane, thus complementing SEGUE and RAVE (which exclude most stars below $\\sim200$~pc), whereas HERMES aims at obtaining chemical information for around $10^6$ stars. In the near future, 4MOST \\citep{dejong12}, probably the most ambitious project, aims at sampling even larger volumes by obtaining chemo-kinematic properties of many millions of stars, taking full advantage of the Gaia astrometric results. The common aim of the huge observational campaigns briefly summarized above is to constrain the MW assembly history -- one of the main goals of the newly emerged field, Galactic Archaeology.  The underlying principle of Galactic Archaeology is that the chemical elements synthesized inside stars, and later ejected back into the interstellar medium (ISM), are incorporated into new generations of stars. As different elements are released into the ISM by stars of different masses and, therefore, on different timescales, stellar abundance ratios provide a cosmic clock, capable of eliciting the past history of star formation and gas accretion of a galaxy\\footnote{In most cases, the stellar surface abundances reflect the composition of the interstellar medium at the time of their birth; this is the reason why stars can be seen as fossil records of the Galaxy evolution.}. One of the most widely used ``chemical-clocks\" is the [$\\alpha$/Fe] ratio\\footnote{Here we use the notation in brackets to indicate abundances relative to the Sun, i.e., [X/Y] $= \\log$(X/Y)$ - \\log$(X/Y)$_{\\odot}$.}.  \\subsection{Galactic Archeology and radial migration}  The power of Galactic Archaeology has been threatened both by observational and theoretical results, showing that stars most probably move away from their birthplaces, i.e, migrate radially. Observational signatures of this radial migration (or mixing) have been reported in the literature since the 1970s, with the pioneering works by \\cite{grenon72, grenon89}. Grenon identified an old population of \\emph {super-metal-rich stars} (hereafter SMR), that are currently located in the solar vicinity, but have kinematics and abundance properties indicative of an origin in the inner Galactic disk (see also \\citealt{castro97} and \\citealt{trevisan11}). These results were extended by \\cite{haywood08}, who showed by re-analyzing the GCS data, that the low- and high-metallicity tails of the thin disk are populated by objects whose orbital properties suggest an origin in the outer and inner Galactic disk, respectively. In particular, the so-called SMR stars show metallicities that exceed the present-day ISM and those of young stars in the solar vicinity. As discussed by \\cite{chiappini03} (see also Table~5 by \\citealt{asplund09}), the metallicity in the solar vicinity is not expected to have increased much since the Sun's formation, i.e., in the last $\\sim$4~Gyr, because of the rather inefficient star formation rate (SFR) at the solar radius during this period, combined with continuous gas infall into the disk. Hence, as summarized in \\cite{chiappini09}, pure chemical evolution models for the MW thin disk cannot explain stars more metal rich than $\\sim$0.2~dex, and radial migration has to be invoked.  N-body simulations have also long shown that radial migration is unavoidable. \\cite{raboud98} studied numerical simulations aimed at explaining the results reported by \\cite{grenon89} of a mean positive U-motion (where U is the Galactocentric radial velocity component of stars), which the authors interpreted as metal-rich stars from the inner Galaxy, wandering in the solar neighborhood. However, \\cite{raboud98} interpreted their findings as stars on hot bar orbits, not recognizing that permanent changes to the stellar angular momenta are possible. It was not until the work by \\cite{sellwood02} that radial migration was established as an important process affecting the entire disk, where stars shift guiding radii due to interaction with transient spiral structure. Modern high-resolution simulations (e.g., \\citealt{roskar08a, roskar12}) have confirmed this finding, but left the role of the Galactic bar unexplored.  A different radial migration mechanism was proposed by \\cite{mf10} and \\cite{minchev11a}, resulting from the nonlinear coupling between the bar and spiral waves, or spirals of different multiplicity \\citep{mq06, minchev12a}. These works, along with studies of diffusion coefficients in barred disks \\citep{brunetti11,shevchenko11}, predict a variation in migration efficiency with time and disk radius, establishing that the dynamical influence of the bar plays an integral part in the MW disk modeling. Recently, \\cite{comparetta12} showed that radial migration can result from short-lived density peaks arising from interference of spiral density waves, even if patterns are long-lived. In addition to internal axisymmetric structure, perturbations caused by minor mergers have also been shown to be effective at mixing the outer disks \\citep{quillen09, bird12}, but can also indirectly affect the entire disk by inducing (reinforcing) spiral and bar instabilities. Considering the established presence of a central bar, spiral structure and evidence for merger activity in the MW, it is clear that all of the above mentioned radial migration mechanisms would have an effect on the Galactic disk.  In summary, a number of both observational and theoretical results challenge the power of Galactic Archaeology. Therefore, the only possible way to advance in this field is the development of chemodynamical models tailored to the MW in the cosmological framework. Only then, a meaningful comparison with the large amounts of current and forthcoming observational data (as summarized in the beginning of this section), can be carried out. This is the main goal of the present work, namely, to develop a chemodynamical model for the MW, to be able to quantify the importance of radial mixing throughout the evolution of our Galaxy.  \\subsection{Difficulties with fully self-consistent simulations} \\label{sec:cos}  Producing disc-dominated galaxies has traditionally been challenging for cosmological models. In early simulations, extreme angular momentum loss during mergers gave birth to galaxies with overly concentrated mass distributions and massive bulges \\citep[e.g.,][]{navarro91,navarro94,abadi03}. Although an increase in resolution and better modeling of star formation and feedback have allowed recent simulations to produce MW-mass galaxies with reduced bulge fractions \\citep{agertz11, guedes11, martig12}, none of these simulations include chemical evolution. Galaxy formation simulations including some treatment of chemical evolution have been performed by a number groups \\citep{raiteri96, mosconi01, lia02, kawata03, kobayashi04, scannapieco05, martinez08, oppenheimer08, wiersma09, few12}. However, although the results are encouraging and globally observed trends seem to be reproduced, such as the mass-metallicity relation \\citep[e.g., ][]{kobayashi07} or the metallicity trends between the different Galactic components \\citep[e.g., ][]{tissera12}, it is still a challenge for these simulations to match the properties of the MW (e.g., the typical metallicities of the different components -- \u00ca\\citealt{tissera12}). Additionally, the fraction of low-metallicity stars are often overestimated \\citep{kobayashi11, calura12}, and reproducing the position of thin- and thick-disk stars in the [O/Fe]-[Fe/H] plane has proved challenging \\citep{brook12}. While these problems might be due to unresolved metal mixing \\citep{wiersma09}, it is also worth noting that none of the above-mentioned simulations reproduces simultaneously the mass, the morphology, and the star formation history (SFH) of the MW.  % This situation has led us to seek a novel approach to address this complex problem. We will show that this approach works encouragingly well, explaining not only current observations, but also leading to a more clear picture regarding the nature of the MW thick disk.  \\subsection{Thick-disk formation scenarios}  The large uncertainties in important observational constraints in the MW, such as the age-velocity-metallicity relation, the abundance gradients and their evolution, together with the inherent complexity of the topic of Galaxy assembly, have led to different scenarios to be proposed for the formation of the thick disk.  One possibility is that thick disks were born thick at high redshift from the internal gravitational instabilities in gas-rich, turbulent, clumpy disks \\citep{bournaud09, forbes12} or in the turbulent phase associated with numerous gas-rich mergers \\citep{brook04,brook05}. They could also have been created through accretion of galaxy satellites \\citep{meza05,abadi03}, where thick-disk stars then have an extragalactic origin.  Another possibility is that thick disks are created through the heating of preexisting thin disks with the help of mergers \\citep{quinn93, villalobos08, dimatteo11}, whose rate decreases with decreasing redshift. Evidence for satellite-disk encounters can be found in structure in the phase-space of MW disk stars (e.g., \\citealt{minchev09, gomez13, gomez12c}), which can last for as long as $\\sim$4~Gyr \\citep{gomez12b}.  Finally, a recently proposed mechanism for the formation of thick disks is radial migration, which we discuss next.  \\subsection{Radial migration and thick disks} \\label{sec:mig}  In the past several years there has been a growing conviction that radial migration, driven by transient spirals, can be responsible for the formation of thick disks by bringing out high-velocity-dispersion stars from the inner disk and the bulge (no need for mergers). This scenario was used, for example, in the analytical model of \\cite{schonrich09a} (SB09a) and \\cite{schonrich09b} (SB09b), where the authors managed to explain the MW thin- and thick-disk characteristics without the need of mergers or any discrete heating processes. Similarly, the increase of disk thickness with time found in the simulation by \\cite{roskar08a} has been attributed to migration in the works by \\cite{sales09} and \\cite{loebman11}.  However, how exactly radial migration affects disk thickening in dynamical models had not been demonstrated (but only assumed) until the recent work by \\cite{minchev11b} and \\cite{minchev12b}. The latter authors showed unambiguously that radial migration driven by secular (or internal) evolution has only a minor effect on disk thickening, mostly beyond three disk scale-lengths, which results in a flared disc\\footnote{\\cite{minchev12b} showed that migrators {\\it contract} (vertically cool) the disk inside the bar's CR, further contributing to disk flaring. This is related to the predominance of inward migrators at such small radii with a mean action smaller than that of the local population.}. This is due to the conservation of vertical action of migrators, as opposed to the incorrect assumption of vertical energy conservation. While not building up a {\\it kinematically} thick disk (since the effect of outward and inward migrators mostly cancels out), \\cite{minchev12b} noted that radial migration in isolated galaxies does contribute to a {\\it chemically} thick disk in the sense that outward migrators are preferentially deposited at higher distances above the galactic plane, with the converse effect for inward migrators. It can be expected that migration {\\it can} thicken the disk, provided that stars were ``preheated\", e.g., either by mergers or by being born hot\\footnote{Note that migration efficiency drops with velocity dispersion \\citep{solway12}.} (both of these are expected at high redshift).  Once again, to advance our understanding of the MW disk formation and evolution, we need to make quantitative estimates of the importance of radial mixing, guided by observational constraints. This implies the need for chemodynamical models making predictions specifically for the solar vicinity, where most of the current observational constraints are found. In Sec.~\\ref{sec:sim} we describe the disk simulation we adopt in the present work. Sec.~\\ref{sec:chem} we describe our input chemistry. Sec.~\\ref{sec:tech} is devoted to our new approach. Our results are shown in Sec.~\\ref{sec:res}, while a new explanation for the origin of the thick disk can be found in Sec.~\\ref{sec:thick}. Here we concentrate on results for the solar vicinity, while results for the whole disk are shown in paper II of this series. Conclusions are drawn in Sec.~\\ref{sec:concl}.   \\begin{figure*} \\includegraphics[width=18cm]{xy.eps} \\caption{ {\\bf First row:} The left panel shows the rotational velocity (dashed blue curve) and circular velocity (solid black curve) at the final simulation time. The middle panel presents the $m=2$ Fourier amplitudes, $A_2/A_0$, as a function of radius estimated from the stellar density. Curves of different colors present the time evolution of $A_2/A_0$. To better see the evolution of the bar strength with time, in the right panel we show the amplitude averaged over the bar maximum. {\\bf Second row:} Face-on density maps of the stellar component for different times, as indicated. {\\bf Third row:} The corresponding edge-on view. Contour spacing is logarithmic. {\\bf Fourth row:} Changes in angular momentum, $\\Delta L$, as a function of radius, estimated in a time window of 0.52~Gyr, centered on the times of the snapshots shown above. Both axes are divided by the circular velocity, thus units are kpc (galactic radius). Strong variations are seen with cosmic time due to satellite perturbations and increase in bar strength. } \\label{fig:xy} \\end{figure*} ",
        "conclusions": "\\label{sec:concl}  In this work we have presented a new approach for studying the chemodynamical evolution of galactic disks, with special emphasis on the Milky Way (MW). Unlike other similar studies, where either the dynamics was too simplistic or the star formation history (SFH) and chemical enrichment was unconstrained, our chemodynamical model is a fusion between a pure chemical evolution model and a high-resolution simulation in the cosmological context. As we argued in Sec.~\\ref{sec:tech}, this new approach allows us to bypass most known problems encountered in fully self-consistent simulations, where chemical enrichment still proves to be a challenge (see Sec.~\\ref{sec:cos}). Moreover, this is the first time that a chemodynamical model has the extra constraint of defining a realistic solar vicinity also in terms of dynamics (see Sec.~\\ref{sec:sim}).  The main results of our chemodynamical model can be summarized as follows:  $\\bullet$ The distribution of birth radii, $r_0$, of stars ending up in a properly defined solar neighborhood-like location after 11.2~Gyr of evolution peaks close to $r_0=6$~kpc due to radial migration (left panel of Fig.~\\ref{fig:r0}). The strongest changes in stellar guiding radii were found for the oldest stars, related to the strong merger activity at high redshift in our simulation and the effect of the bar at later times. Locally born stars of all ages can be found in the solar neighborhood. While a wide range of birth radii is seen for different age groups, the majority of the youngest stars are born at, or close to, the solar neighborhood bin.  $\\bullet$ While the low-end in our simulated solar neighborhood metallicity distribution function (MDF) is composed of stars with a wide range of birth radii, the tail at higher metallicities ($0.25<$[Fe/H]$<0.6$) results almost exclusively from stars with $3<r_0<5$~kpc (see Fig.~\\ref{fig:r0}, middle panel). This is the region just inside the bar's CR, which is where the strongest outward radial migration occurs, as we discussed in Sec.~\\ref{sec:sim} and showed in the bottom row of Fig.~\\ref{fig:xy} (see also Fig.~\\ref{fig:ap1} and associated discussion). The fraction of stars in this tail can, therefore, be related to the bar's dynamical properties, such as its strength, pattern speed, and time evolution/formation.  $\\bullet$ Constraining our simulated sample spatially, we provided excellent matches to the metallicity distribution of stars confined close to the Sun (data from \\citealt{adibekyan12}) and the SEGUE G-dwarfs (\\citealt{bovy12b} and Brauer et al. 2013), which covers approximately $0.5<|z|<3$~kpc, predicting the expected shift in the distribution peak from solar (nearby sample) to [Fe/H]~$\\sim-0.3$ for the SEGUE coverage.  $\\bullet$ Our results suggest that the most likely birth location for the Sun is in the range $4.4<r_0<7.7$~kpc, with the highest probability $\\sim$5.6~kpc, followed by $\\sim$7~kpc (Fig.~\\ref{fig:r0}, right). This estimate comes from both dynamical and chemical constraints and is dependent on the migration efficiency in our simulation and the adopted chemical evolution model.  $\\bullet$ Examining the effect on the age-metallicity relation (AMR), we found that some flattening is observed, mostly for ages $\\gtrsim9$~Gyr (Fig.~\\ref{fig:chem}). Although significant radial mixing is present, the slope in the AMR is only weakly affected. This is related to the intermediate distance of the Sun from the Galactic center, where the effect of stars arriving from the inner and outer disks is well balanced, while the migration efficiency is exposed by the scatter around the mean. Therefore, large scatter does not necessarily imply flattening.  $\\bullet$ We found no bimodality in the [Fe/H]-[O/Fe] stellar density distribution. However, when selecting particles according to kinematic criteria used in high-resolution samples to define thin and thick disks, we recover the observed discontinuity (Fig.~\\ref{fig:o_fe1}). This agrees with the recent observational results by \\cite{bovy12a}, where a smooth [Fe/H]-[O/Fe] distribution was obtained after correcting for the spectroscopic sampling of stellar subpopulations in the SEGUE survey. A very good match was achieved between our model and the recent high-resolution data by \\cite{ramirez13}.  $\\bullet$ The vertical mass density in our simulated solar vicinity is well fit by the sum of two exponentials giving thin- and thick-disk scale-heights $h_1=330$~pc and $h_2=1200$~pc, respectively (Fig.~\\ref{fig:den}, top). By separating this sample into narrow bins of age, [O/Fe], and [Fe/H], we found that the vertical scale-height of each component can be fitted well by a single exponential, with values growing with increasing age, increasing [O/Fe], and decreasing metallicity. An abrupt increase in scale-height was found for samples with age~$\\gtrsim8$~Gyr, [O/Fe]~$\\gtrsim0.15$, and [Fe/H]~$\\lesssim0.5$, related to stars born with high velocity dispersions and strong merger activity at high redshift. The vertical velocity dispersions for each of these subpopulation was found to exhibit smooth variations with height above the disk plane, strongly flattening for old high-[O/Fe] metal-poor stars (Fig.~\\ref{fig:den}, bottom).  $\\bullet$ A good match was achieved between the mean metallicity variation with height above the disk plane resulting from our model and the data by \\cite{schlesinger12} (Fig.~\\ref{fig:fe_z}). A good agreement between our model prediction and the model by \\cite{bovy12a,bovy12b} was found for the metallicity variation with $|z|$, including a flattening at $|z|\\gtrsim1.5$~kpc. This was associated with the predominance of old/high-[O/Fe] stars at large distances from the plane (see Fig.~\\ref{fig:fe_vz}). The metallicity gradients in the $|z|-$[Fe/H] and $|v_z|-$[Fe/H] relations mostly disappear for stars grouped in narrow age- or [O/Fe]-bins (Fig.~\\ref{fig:fe_vz}). This provides additional predictions for Galactic surveys.  $\\bullet$ For stars in the range $3<$~age~$<8$~Gyr, our model predicts a mostly flat age-velocity relation (AVR) for the radial component, $\\sigma_r$, and a weak increase for the vertical component, $\\sigma_z$ (Fig.~\\ref{fig:sig}). For ages older than 8~Gyr a continuous strong increase is observed for both components, which is related to the massive mergers and stars born with high velocity dispersions at this early epoch. These hot stars can be associated with the MW thick disk. The vertical velocity dispersion of the oldest stars in the absence of mergers (and/or stars born with high velocity dispersions) is underestimated by a factor of $\\sim2$, suggesting a violent origin for the MW thick disk (see appendix A, Fig.~\\ref{fig:ap1} bottom row).  $\\bullet$ We found a strong flattening in the [Fe/H] radial profiles for the older populations with the effect diminishing strongly for the younger ones (see Fig.~\\ref{fig:grad}). Stars younger than 2~Gyr have a final gradient very similar to the initial one out to $\\sim12$~kpc, justifying its use as a constraint for our chemical model.  $\\bullet$ We predict that the [O/Fe] radial profiles are essentially preserved for the chemical model we used. The [O/Fe] profiles for different age groups result straightforwardly from the adopted variation of the infall-law with radius (and hence the SFHs at different positions) and, thus, provide a way to constrain different chemical evolution models. In the near future, these will be able to be measured by combining the good distances and ages expected from the CoRoT mission \\citep{baldin06}, with abundance ratios obtained by spectroscopic follow-up surveys. For the young populations, this should be already possible to obtain from the observations of open clusters, e.g., with the ongoing Gaia-ESO or APOGEE surveys.  $\\bullet$ Finally, probably one of the most important outcomes of our chemodynamical model is that although we used only a thin-disk chemical evolution model, the oldest stars that are now in the solar vicinity show several of the properties usually attributed to the Galactic thick disk. According to the results of the present work, the MW thick disc has emerged naturally from stars migrating from the inner disk very early on due to strong merger activity, followed by additional radial migration driven by the bar and spirals at later times (see Sec.~\\ref{sec:liu}). Importantly, a significant fraction of old stars with thick-disk characteristics could have been born near to and beyond the solar radius (Fig.~\\ref{fig:liu}, middle row).  We showed that hot old metal-poor high-[$\\alpha$/Fe] stars can be delivered to the solar neighborhood at high redshift as the effect of massive perturbers, such as the one with 1:5 mass ratio described in this work (see Fig.~\\ref{fig:liu}). For a MW-size galaxy, the occurrence of mergers of this size around $z\\sim1$ is consistent with the numerical results by a number of groups (e.g., \\citealt{benson04, stoehr06, kazantzidis08, delucia08, stewart08, villalobos08, purcell09}).  Alternatively, a bar can also bring some fraction of these stars to the solar neighborhood, but over a long period of time, as we showed in Fig.~\\ref{fig:liu}, bottom row. However, a difference in the metallicity and velocity distributions of these old stars would result. Most notably, avoiding a merger not only avoids the extreme migration at high redshift, but also does not result in high enough velocity dispersions today. For example, we found AVRs with maxima at $\\sigma_r\\approx50$~km/s and $\\sigma_z\\approx28$~km/s in the lack of strong mergers, to be compared to $\\sigma_r\\approx65$ and $\\sigma_z\\approx55$~km/s for our standard model. Additionally, the ratio of cold to hot orbits (samples separated by an eccentricity cut at 0.2) changes from $\\sim$0.33 to $\\sim0.96$ when the merger is avoided. This large difference can be used as a constraint in properly mass-corrected observational data to assess the possibility of an early-on merger in the MW.  It is therefore tempting to conclude that the highest velocity dispersion stars observed in the solar neighborhood are a clear indication of an early-on merger-dominated epoch, or equivalently, scenarios involving stars born with high velocity dispersions. Our results suggest that the MW thick disk cannot be explained by a merger-free disk evolution as proposed previously. This is related to the inability of internally driven radial migration to thicken galactic disks \\citep{minchev12b}. However, an observational test involving both chemical and kinematic information must be devised to ascertain the effect of early mergers.  While here we concentrated on the solar vicinity, results for the entire disk will be presented in paper II of this series. Chemodynamical predictions for different disk radii may help further constrain thick-disk formation scenarios and the MW assembly in general.  The results of this work present an improvement over previous models as we used a state-of-the-art simulation of a disk formation to extract self-consistent dynamics and fuse them with a chemical model tailored for the MW. Future work should consider implementing gas flows, Galactic winds, and Galactic fountains in the model to asses the importance of these processes. A suite of different chemical evolution models assigned to different disk dynamics (from different simulations) should be investigated to make progress in the field of Galactic Archeology. Another next step is providing model predictions for the MW disk properties at different Galactocentric regions and for a large number of chemical elements (especially in view of the ongoing high-resolution surveys, such as HERMES and APOGEE).  The method presented here is not only applicable to the MW, but is also potentially very useful for extragalactic surveys. The same approach can be used for any other galaxy for which both kinematic and chemical information is available. A large such sample is currently becoming available from the ongoing CALIFA survey \\citep{sanchez12}. Given the SFH and morphology of each galaxy, a chemodynamical evolution model tailored to each individual object can be developed, thus providing a state-of-the-art database for exploring the formation and evolution of galactic disks."
    },
    "1208/1208.0886_arXiv.txt": {
        "abstract": "We compare calculated intensities of lines of \\cii, \\ni, \\nii\\, \\oi\\ and \\oii\\ with a published deep spectroscopic survey of \\ic. Our calculations use a self--consistent nebular model and a synthetic spectrum of the central star atmosphere to take into account line excitation by continuum fluorescence and electron recombination. We found that the \\nii\\ spectrum of s, p and most d states is excited by fluorescence due to the low excitation conditions of the nebula. Many \\cii\\ and \\oii\\ lines have significant amounts of excitation by fluorescence. Recombination excites all the lines from f and g states and most \\oii\\ lines. In the neutral--ionized boundary the \\ni\\ quartet and \\oi\\ triplet dipole allowed lines are excited by fluorescence, while the quintet \\oi\\ lines are excited by recombination. Electron excitation produces the forbidden optical lines of \\oi\\, and continuum fluorescence enhances the \\ni\\ forbidden line intensities. Lines excited by fluorescence of light below the Lyman limit thus suggest a new diagnostic to explore the inner boundary of the photodissociation region of the nebula. ",
        "introduction": "With an apparent simple geometry the low excitation planetary nebula \\ic\\ lends itself to simpler models with fewer assumptions on its structure and kinematics than are generally necessary to reproduce observations of planetary nebulae (PNe). A recent deep spectroscopic survey of this object by \\citet{emili} (hereafter SWBH) identified several hundred lines in \\ic\\ from which \\citet*{sharpee} (hereafter SBW) derived ionic abundances for CNONe ions relative to H. Many dipole--allowed line intensities from those ions in that survey cannot be explained by the recombination theory with the same abundances derived from collisionally excited lines (CELs). The intensities of recombination lines are often used to measure ionic abundances because they depend more weakly on temperature and density than CELs. Abundances derived from recombination lines, however, are usually higher than those derived from CELs in the Orion nebula \\citep*{esteban,baldwin00,esteban04,mesa}, many \\hii\\ regions \\citep{garcia07}, and planetary nebul\\ae\\ \\citep*{liu95,liu01b,luo,liu06,tsamis08,garcia09,otsuka}. \\citet{williams} showed that CELs probably give the correct abundances because the values derived from CELs are consistent with those derived from UV absorption lines in a sample of four PNe. The discrepancies of less than a factor of 2 found by SBW in \\ic\\ between abundances from CELs and dipole--allowed lines are smaller than in other objects. Nevertheless it is important to ensure that the intensities of the dipole--allowed lines are not enhanced by additional excitation mechanisms besides recombination whenever they are used in temperature or density diagnostics \\citep*{fang} (hereafter FSL) or abundance determinations.  The usual procedure to determine a species X abundance directly from optical line intensities is to measure the ratio of the intensity of a line emitted by the species with respect to the \\hb\\ intensity and assume that the ratio of column densities producing the lines is equal to the abundance. Specifically, the abundance X/H is taken from the equation \\beq {I(\\lambda)\\over I(\\hb)}={\\int\\epsilon(\\lambda)\\,dV \\over\\int\\epsilon(\\hb)\\,dV}= {\\langle\\epsilon(\\lambda)\\rangle\\over\\langle\\epsilon(\\hb)\\rangle} {\\rm X\\over H} \\label{ionabun} \\eeq where $I(\\lambda)/I(\\hb)$ is the ratio of intensities corrected for reddening, $\\epsilon(\\lambda)$ and $\\epsilon(\\hb)$ are the emissivities of a line emitted by X and of \\hb\\ respectively, both are integrated over some volume fraction of the nebula covered by the aperture, and $\\langle\\epsilon\\rangle$ means some kind of average of the emissivities over that volume fraction. Averaged emissivities over the volume within the solid angle of the aperture imply the adoption of average values for the temperature, density and filling factors.  An alternative procedure to determine chemical abundances in a nebula is to propose mutually consistent models of the nebula and the atmosphere of the exciting star or stars. The stellar parameters, chemical abundances and density profile of the nebula are inputs to a photoionization model, and their values are fixed by matching the predicted line intensities to the observations. The variation of stellar and nebular parameters of the model can be constrained by direct observations of the star and nebula to reduce the degeneracy of different parameters giving similar matches to the observed line intensities. \\citet{chris09} (hereafter MG) showed that a consistent stellar energy distribution (SED) of the central star of \\ic\\ and a photoionization model give an impressive agreement of predictions and measured intensities of CELs and many recombination lines within observational errors. However most dipole--allowed lines in the SWBH survey were not included in the model by MG, and some of the ones that were included show observed intensities much larger than their predicted values because the MG calculations did not include fluorescence excitation.  In this work we add fluorescence excitation of permitted lines to a photoionization model and a SED similar to the ones used by MG, and compare the results with all lines observed by SWBH that have probable identifications as dipole--allowed lines of \\cii, \\ni, \\nii, \\oi, and \\oii\\ and optical forbidden lines of \\ni\\ and \\oi. Continuum fluorescence of starlight was suggested as a mechanism to excite permitted lines in ionized nebulae by \\citet{seaton68} and \\citet{grandi75,grandi76}. \\citet{orion} (hereafter EM) showed that fluorescence in Orion dominates the intensities of \\nii\\ permitted lines. The low excitation level of \\ic\\ and the resemblance of its ionization degree to that of the Orion nebula noted by \\citet*{torres} suggest that a similar fluorescence mechanism may occur in both objects. ",
        "conclusions": "We have calculated the intensities of weak emission lines of \\cii, \\ni, \\nii, \\oi\\ and \\oii\\ using a consistent nebular model and SED, which proved successful in previous work in reproducing the intense emission lines and overall density diagnostics and surface brightness maps of \\ic. We compared our calculations with a published deep spectroscopic survey taking into account the fine structure of the lines and including some weak lines with possible blends. Although this nebula shows an almost spherical symmetry, our calculations do take into account the position, size and orientation of the aperture used in the observations.  As in our previous work on the Orion nebula, we found that fluorescence is responsible for the excitation of lines from s, p and d states in \\nii. Fluorescence also contributes significantly to the excitation of most lines of the s, p and d states in \\cii\\ and some p and d states of \\oii\\ while the rest of the \\oii\\ states are mostly excited by recombination. Some lines excited by fluorescence are sensitive diagnostics of the details of the SED, the expansion velocity gradient and turbulence of the nebula, the transition probabilities involved in the pumping mechanism, and their intensity is closely related to lines in the stellar atmosphere. The recombination rates can explain most of the observed intensities from f and g states of all species in the ionized region although there are many observed transitions with excited core configurations awaiting the calculation of their effective recombination coefficients.  Ionic abundances deduced from dipole--allowed lines assuming pure recombination rates without taking into account fluorescence excitation of the lines can be overestimated by factors of up to 10 for \\cii, 200 for \\nii, 2 for \\oii, and 500 for \\ni\\ and \\oi\\ depending on the transition being considered, and a detailed nebular model and SED are needed to properly estimate ionic abundances from lines mostly excited by fluorescence.  The automated procedure of line identification with the \\emi\\ code seems to be in accordance with our calculations because the largest discrepancies between predicted and observed line intensities tend to occur in lines with dubious identifications or blends as indicated by that code. A deep spectroscopic nebular survey coupled with an appropriate line identification procedure can prove useful to test the accuracy of quantum mechanical calculations of atomic data.  The \\ni\\ quartet and \\oi\\ triplet lines are produced by fluorescence in a narrow and dense interface of the neutral and ionized regions. Electron collisions compete with fluorescence in the excitation of the [\\ni]\\lam5200.26 and \\lam5197.90 lines while fluorescence contributes a negligible amount to the excitation of the [\\oi]\\lam6300.30, 6363.78 and \\lam5577.34 lines. However our calculations tend to underestimate systematically the observed intensities of the [\\ni] lines by a factor of 5, and further work on the modeling of the PDR is necessary to fully understand the formation of those lines."
    },
    "1208/1208.2284_arXiv.txt": {
        "abstract": "{A sample of 229618 narrow emission-line galaxies is used to establish two new unambiguous type of evidence for supermassive black holes at the center of their nuclei: 1) the Seyfert 2 galaxies and LINERs follow the same characteristic power law relating the luminosity of ionized flux with that of the continuum; 2)  both show the highest concentration of mass at their center, independent of the morphology of the galaxy, consistent with  higher binding energies. The Full Width at Half Maximum is shown to be related with the mass concentration, suggesting that the kinetic energy of the gas in AGNs has a gravitational origin.  Within the standard accretion model, the Transition-type Objects, Seyfert 2 galaxies and LINERs represent AGNs forming supermassive black holes on different mass-scales, or they could be related through an evolutionary process, the LINERs representing the end product of this evolution. }  \\resumen{Una muestra de 229618 galaxias con l\\'ineas angostas de emisi\\'on es utilizada para establecer dos nuevos tipos de evidencias inambiguas de agujeros negros supermasivos en el centro de sus n\\'ucleos: 1) las galaxias Seyfert 2 y los LINERs siguen la misma ley de potencia caracter\\'istica relacionando la luminosidad del flujo ionizante con la del continuo; 2) ambas muestran la m\u00e1s alta concentraci\\'on de masa en sus centro, independientemente de la morfolog\\'ia de la galaxia, consistente con altas energ\\'ias de ligaci\\'on. El ancho a potencia media est\\'a relacionado con la concentraci\\'on de masa, sugiriendo que la energ\\'ia cin\\'etica del gas en los AGNs tiene un origen gravitacional. Dentro del modelo est\\'andar de acreci\\'on, los Objetos de Transici\\'on, las galaxias Seyfert 2 y los LINERs representan AGNs formando agujeros negros supermasivos a diferentes escalas en masa, o prodr\\'ian estar relacionados dentro de un proceso de evoluci\\'on, donde los LINERs ser\\'{\\i}an el producto final. }   \\addkeyword{galaxies: activity} \\addkeyword{(galaxies:) quasars: general} \\addkeyword{galaxies: morphology} \\addkeyword{galaxies: accretion rate}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  In the early 1960's, numerous unresolved radio sources were found that are now known as quasars \\citep{Weedman86}. The optical spectra of quasars show they are extremely rich in hot gas, ionized by copious amounts of high-energy photons \\citep{Osterbrock89,BNW90,Krolik99}. Very rapidly it was discovered that there are much more radio-quiet quasars than radio-loud ones \\citep{Sandage65,Braccesi80,Schmidt83,Padovani93,deVries06}. Understanding and explaining the optical spectra of quasars was also instrumental in confirming the interpretation of the redshift in terms of the expansion of the universe \\citep{Schmidt63,Heckman84}. Adopting this cosmology and putting all the observations together, with time it became clear that the energy produced by a quasar is so phenomenal that the only physical mechanism known that could explain it is the transformation of gravitational energy into radiation through the accretion of matter on a supermassive black hole (SMBH) forming at the center of galaxies \\citep{Rees78,Soltan82,Netzer85}. Today, this interpretation constitutes the basis of the standard model of the Active Galactic Nuclei (AGN) phenomenon.  Subsequent observations in the 1970's and 1980's rapidly showed that quasars were much more common at high redshift, that is, in the past, which is now largely accepted as evidence of luminosity or/and density evolution with time \\citep{Schmidt72,Braccesi80,Veron86,Yee87,Boyle00,Richards06,Croom09}. The common interpretation for this evolution is that luminous AGNs mark the first phase in the formation of galaxies by mergers, possibly following, very tightly and in a complex intricate way, the formation of their bulges \\citep{Cavaliere86,Silk98,Miller03,DiMatteo05,Lapi06,Hopkins07a,Horst07,Elbaz09,Letawe10,Treister10}. Consistent with this interpretation it was found that the mass of the SMBH is well correlated with the bulge mass or the stellar velocity dispersion of the galaxies \\citep{Magorrian98,Ferrarese00,Gebhardt00,Haring04,Peterson05,Gultekin09,Graham11}.  So far, the standard model for AGN seems to offer an elegant and credible explanation for quasars, in good agreement with cosmology. However, here lies one mystery: if one counts the number of luminous AGNs at high redshift, then almost all massive galaxies in the nearby universe must have had a comparable high-activity phase during their formation \\citep{Kormendy95,Richstone98}. But what is the evidence at low redshift in favor of such a conclusion? In other words, where are the luminous quasars, or rather their remnants, at low redshift \\citep{Lopez-Corredoira11}? Within this interpretation, the luminous AGN phase is assumed to be very short in duration ($10^6$ to $10^8$ yrs) compared to the formation and evolution of galaxies \\citep{Martini04,Kelly10}. This suggests that most quasar remnants today should not show any trace of non-stellar activity in their nucleus. But how short really is this phase \\citep{Cattaneo01}, and how consistent is this scenario with the low level of nuclear activity observed in galaxy surveys at low redshift \\citep{Ueda03,Merloni04,Heckman04,Barger05}?  In the present study we use one of the largest spectroscopic survey to date, the Sloan Digital Sky Survey \\citep[SDSS;][]{York00,Stoughton02}, in order to shed some light on the evolution of AGNs at low luminosity. The SDSS spectral survey has revealed that a high fraction of nearby galaxies possess ionized gas in their nucleus. A small fraction of these galaxies, the Seyfert~1 (Sy1), show broad emission-line components akin to what is observed in quasars \\citep{Weedman86,Dibai87,Osterbrock89,Krolik99}. The standard model of AGN suggests that the Sy1s are similar to quasars, forming the low-luminosity tail of their distribution \\citep{Smith86,Marshall87}. In terms of the SMBH theory it is expected that these AGNs have lower-mass black holes for the same accretion luminosity, $L_{acc}$, and efficiency, $\\eta$, than the more luminous AGNs \\citep{Hopkins06a}.  However,  much more numerous in the SDSS survey are narrow emis\\-sion-line galaxies (NELGs). Using diagnostic diagrams that compare various emis\\-sion-line ratios \\citep{Baldwin81,Veilleux87}, and different classification criteria \\citep{Kewley01,Kauffmann03}, a good fraction of these NELGs were found to have spectra consistent with that of an AGN, although they do not show as broad lines as the Sy1 and are less luminous. One of these AGNs, the Seyfert~2 (Sy2), is thought to differ from the Sy1s only because the broad-line regions in these galaxies are obscured by a circumnuclear torus of gas and dust \\citep[this is the so-called unified AGN model;][]{Antonucci85}. On the other hand, the Sy2s may also be intrinsically different from the Sy1s, lacking, for some unknown reason, the broad-line regions.  Another type of AGN which is even more problematic than Sy2 is the Low Ionization Nuclear Emission-line Regions \\citep[LINER;][]{Heckman80,Coziol96,Kewley06}. Despite the classification based on standard spectral diagnostic diagrams, which identify these galaxies as AGNs similar to Sy2s, many researchers today claim that LINERs are ``false AGNs'', mostly because many show active and intense star-formation activity in their circumnuclear regions, and do not present clear evidence for a compact source in their nucleus. The presence of star formation in narrow-line AGNs seems now indubitable \\citep{Kauffmann03}. This is why the Transition-type Object (TO) class was systematically included, defining in diganostic diagrams a sort of ``buffer'' zone between pure Star Forming Galaxies (SFGs) and pure (or rather dominant) AGNs. The presumed absence of a compact source in LINERs, on the other hand, was never proven.  Instead, evidence for an AGN contribution is usually downgraded in favor of star formation, but never eliminated completely \\citep[e.g.,][]{Annibali10}.  Finding evidence of SMBHs genuine AGNs should not pose such a problem. For example, in one of his books, \\citet{Osterbrock89} proposed a simple test, which consists in comparing the luminosity of the emission line, say H$\\alpha$, with the continuum luminosity in the blue, say at 4800\\AA\\ (sufficiently far from H$\\beta$ at 4861\\AA). According to the standard model of AGN, the source of the ionizing flux is the UV continuum produced by the accretion of matter onto a SMBH. Consequently, a typical AGN must show a characteristic power law relation, $L_{em}\\propto L^{\\alpha}_{cont}$, significantly different from any relation followed by SFGs (the same principle, for the same reasons, is at the basis of the spectral diagnostic diagrams). Up to now, however, this test was never applied to a large and statistically significant sample of NELGs \\citep[e.g.,][]{Buson93,Macchetto96}. This is one of the major goals of our study.  The plan of this article is the following. In Section~2, we describe the selection of our sample of NELGs and explain how the data for our analysis were obtained. In Section~3 we present the results of our spectral classification analysis. In Section~4 we study and discuss the results of the Osterbrock power-law test and in Section~5 we present other evidence consistent with SMBHs at the center of NELGs. Finally, in Section~6 we discuss how the standard accretion model can explain the different types of AGNs. Our conclusions can be found in Section~7. ",
        "conclusions": "We have shown that the NELGs, identified as LINERs and Sy2s, trace similar power laws between the luminosity in emission and the luminosity in the continuum, the LINERs differing from the Sy2s only by their lower luminosity. This result is consistent with the standard model of AGN, which states that the principal source of ionizing photons in AGNs is the accretion of matter onto a SMBH at their center. Also consistent with the presence of a SMBH at the center of NELGs, we have shown that the VMA, that is, the mass concentrated at the center of the galaxies, is always higher in the AGNs than in galaxies showing a higher level of star formation, like the TOs and SFGs. The fact that this is independent of the morphology of the galaxies suggests this characteristic is not merely related to the formation of the bulge, but to a more fundamental aspect of the formation of the galaxies. This higher concentration of mass in the AGNs is related with a higher gravitational binding energy due to the formation of a compact massive object, a SMBH, at their center.  We have also found a strong correlation between the VMA and the FWHM, which suggests that the energy necessary to produce the large movement of gas in the AGNs is gravitational in nature, related with the binding energy of the SMBH at their center.  We conclude that the standard AGN model depicts in an extremely consistent way the narrow-line AGN phenomenon observed at low redshift. According to this model SMBHs form at the center of most galaxies, growing, through active star formation, with the mass of their bulges \\citep[see][and references therein]{Cattaneo01}. Therefore, one should not be surprised to find so ample evidence of AGNs with star formation. The TOs are possibly one good example. As the mass of the bulge and of the SMBH grow the galaxy transforms into an earlier morphological type. Eventually star formation fades away, leaving only a narrow line AGN (a LINER), where a SMBH still accretes matter but at a lower rate than before. According to our findings, it is very possible that all the narrow-line AGNs passed by a phase of higher activity in the past, which could suggest that most LINERs were Sy2s sometime in the past, and most probably, the Sy2s will eventually end their life as LINERs.  The standard AGN model explains very well the $\\nu$-shaped distribution traced by the NELGs in the BPT-VO diagram. This distribution reflects the  accumulation of mass at the center of galaxies through the formation of a SMBH and a massive bulge. On the left branch of the $\\nu$-shaped distribution, we find galaxies with small binding energies. These galaxies form most of their stars in a disk and do not develop massive bulges.  The relatively small-mass SMBH that form in their nucleus do not accrete much matter either, which explains why we do not detect them. Our own galaxy is possibly of this kind. On the right branch of the $\\nu$-shaped distribution, going from left to right, we find galaxies with increasing binding energies, and consequently more massive bulges and actively accreting SMBHs. The TOs, Sy2s and LINERs in this scenario may be related to different mass scales in the formation of a SMBH. Alternatively they may represent different evolutionary phases which would explain why the exact boundaries between the activity types are difficult to delineate. The LINERs in particular may represent the end products of this evolutionary process.  Within the standard AGN model interpretation the Sy2s and LINERs seem to show clear evidence of evolution (within each type): as the SMBH grows in mass its accretion rate decreases.  This would suggest that the activity phase of the SMBH in narrow-line AGNs is of longer duration than usually assumed based on quasar models, varying over a period of time of the order of a Gyr or more \\citep[e.g.,][]{Kewley06},  larger than the time necessary to form their bulges. In fact, in the narrow-line AGNs we may still see SMBHs growing in mass by accretion, in parallel with the host galaxies building their bulges through star formation (the Sy2s, and possibly also the TOs, are good examples). This longer activity timescale however may not apply to the more luminous AGNs like the quasars, suggesting possibly different accretion regimes \\citep[e.g.,][]{SmallBlandford92,Hopkins06b}. In particular, we see no evidence of negative feedback from the activity of the SMBH, as suggested to regulate the activity of the quasars \\citep[see][and reference therein]{Page12}. This model would not explain the variation of bulge characteristics observed in the narrow-line AGNs, passing from the TOs to the Sy2s and LINERs, which is more consistent with a positive feedback."
    },
    "1208/1208.4557_arXiv.txt": {
        "abstract": "{It has long been suggested that circumstellar disks surrounding young stars may be the signposts of planets, and still more since the recent discoveries of embedded substellar companions. The planet-disk interaction may create, according to models, large structures, gaps, rings or spirals, in the disk. In that sense, the Herbig star HD\\,142527 is particularly compelling as, its massive disk displays intriguing asymmetries that suggest the existence of a dynamical peturber of unknown nature. } {Our goal was to obtain deep thermal images of the close circumstellar environment of HD\\,142527 to re-image the reported close-in structures (cavity, spiral arms) of the disk and to search for stellar and substellar companions that could be connected to their presence.} {We obtained high contrast images with the NaCo adaptive optics system at the Very Large Telescope in L$~\\!'$-band. We applied different analysis strategies to probe the presence of extended structures or point-like sources using both classical PSF-subtraction and angular differential imaging.} {The circumstellar environment of HD\\,142527 is revealed at an unprecedented spatial resolution down to the sub arcsecond level for the first time at $3.8~\\mu m$. Our images reveal important radial and azimuthal asymmetries which invalidate an elliptical shape for the disk as previously proposed. It rather suggests a bright inhomogeneous spiral arm plus various fainter spiral arms. We also confirm an inner cavity down to $30$~AU and two important dips at position angles of $0$ and $135$~deg. The detection performance in angular differential imaging enables the exploration of the planetary mass regime for projected physical separations as close as $40$~AU. The use of our detection map together with Monte Carlo simulations sets stringent constraints on the presence of planetary mass, brown dwarf or stellar companions as a function of the semi-major axis. They severely constrain the presence of massive giant planets with semi-major axis beyond $50$~AU, i.e. probably within the large disk's cavity that radially extends up to $145$~AU or even further outside.} {} ",
        "introduction": "Since the discovery of the first exoplanet \\citep{mayor95} around the solar-like star 51 Pegasi, more than 760 have been detected so far using different observing techniques, from radial velocity measurements, photometric transit, microlensing to direct imaging. These discoveries reveal a great diversity of planetary systems in terms of physical properties (atmosphere and interior) and environment (single or multiple planets around single stars, in a circumstellar disk or even circumbinary planets). A new paradigm about the formation, structure and composition of planets is emerging, wider than what we have learned so far from the Solar System. Moreover, most of the known detected giant planets orbit within $5$~AU owing to the detection biases of radial velocity and transit techniques. The statistical properties of this population of close-in planets, the period-mass distribution, the planet-metallicity correlation together with the densities of Hot Jupiter interiors \\citep{udry07, mayor11} favor a formation mechanism by core accretion followed by gas capture \\citep[CA hereafter,][]{pollack96}.  At wider orbits, the distribution of giant planets is less well known. However, the discoveries of planetary mass companions to young, nearby stars in direct imaging (AB~Pic~b, \\citealt{chauvin05}; RXJ\\,1609~b, \\citealt{lafreniere08}), and further more around Fomalhaut \\citep{kalas08} and HR8799 \\citep{marois08, marois10} challenge the CA model. This scenario suffers from exceedingly long timescales and low disk surface density at large separations. Alternative scenarii could operate at wide orbits, either combined with CA such as planet-planet scattering \\citep{crida09} and outward planetary migration \\citep[see for instance][] {morda12} or alternative formation processes must be invoked like gravitational instability \\citep[GI hereafter][]{cameron78} or stellar binary mechanisms. These formation processes should lead to different physical and orbital distributions of giant planets \\citep{boley09, dodson09}. Hence, new discoveries are needed to understand planet formation at all orbits.  The technique of direct imaging technique also offers the possibility to detect giant planets in circumstellar disks to directly study the planet-disk interaction and even dynamically constrain the planets's mass \\citep{chiang09, lagrange12a}. Indirect signs of planetary formation such as clumps, gaps, holes and spiral arms \\citep[see a review in][]{papaloizou07} have been probably already observed within several proto-planetary disks (e.g AB Aur, \\citealt{fukagawa04}; TCha, \\citealt{huelamo11};  \\citealt{andrews11}SAO~206462, \\citealt{muto12}).  In this context, using high contrast imaging, we have observed the close environment of the star HD\\,142527, known to host an extended proto-planetary disk.  HD\\,142527 is a Herbig star \\citet{waelkens96}, classified as a F6IIIe \\citep{houk78}. The Hipparcos measurement has been revisited by \\citet{leeuwen07} leading to a parallax of $\\pi=4.29\\pm0.98$~mas, i.e. $d=233^{+69}_{-43}$~pc. However, HD\\,142527 has been identified as a member of the star-forming region Sco OB2-2 \\citep{acke04} or Upper Centaurus Lupus \\citep{zeeuw99} hence the distance would be $d=145\\pm15$~pc in both cases. We choose to adopt this latter distance since we consider the identification of the membership more reliable. The age and the mass of the star are also matter of debate since \\citet[F06 hereafter]{fukagawa06} derived an age of $2.2^{+3}_{-2}$~Myr and $M=1.9\\pm0.3$~M$_\\odot$ from its stellar UV/optical luminosities and isochrone \\citep[from][]{palla99} whereas \\citet{verhoeff11} found an age of $\\simeq5$~Myrs and $M=2.2\\pm0.3$~M$_\\odot$ based on a new estimation of the stellar luminosity and pre main sequence isochrones \\citep[from][]{siess00}. Finally, \\citet{pecaut10} discussed the age of F-type members of the Sco-Cen OB association and estimated it to be in the range $[5;20]~$Myr. We therefore considered an age between $2$ and $20~$Myr for further analysis. The Table~\\ref{tab:star_prop} summarizes the stellar properties of HD\\,142527. This system, resolved by F06 in NIR and \\citet{fujiwara06} in IR, presents a circumstellar disk quite unusual with unprecedented infrared excess ($F_{IR}/F_\\star=0.98$ and $F_{NIR}/F_\\star=0.32$, \\citealt{dominik03}) compared to other Herbig stars. Evidence of grain growth (deduced from thermal and scattered light observations) together with the high level of crystallinity (MIDI observations, \\citealt{boekel04}, and NIR spectrum, \\citealt{honda09}) reveal that the outer disk is characterized by  highly-processed dust. Moreover, F06 notes the presence of a spiral arm, a prominent cavity up to $\\simeq150$~AU, two asymmetric facing arcs within the disk and measured a stellar offset of $\\simeq 20$~AU with respect to the center of the disk. Models have been computed to fit the infrared excess, the geometry of the inner disk and the mineralogy of the dust \\citep{verhoeff11}. They found a disk cavity between $30$ and $130$~AU separating an inner disk embedded in a halo and a massive outer disk extended to $200$~AU. They conclude that on-going giant planet formation is suggested from the highly evolved status of the disk surrounding HD\\,142527.  To get a more accurate view of the close-in structures (cavity, spiral arms) of the disk at unprecedented spatial resolution and to search for stellar and substellar companions that could be connected to their formation, we carried out VLT/NaCo observations at L$~\\!'$ band using angular differential imaging (ADI hereafter). We present the observing strategy and the data reduction in Section \\ref{sec:obs}, the fine structures and the geometry of the cavity and of the outer disk in Section \\ref{sec:disk}, finally the detection performance and the constraints on the presence of giant planets in Section \\ref{sec:planets}.  \\begin{table}[th] \\caption{\\label{tab:star_prop}HD\\,142527 properties} \\centering \\begin{tabular}{lc} \\hline\\hline\\noalign{\\smallskip} Parameters & \\\\ \\hline\\hline\\noalign{\\smallskip} RA (J2000)  & $+15:56:41.88$ \\\\ DEC (J200) & $-42:19:23.27$ \\\\ Proper motion RA (mas/yr) & $-11.2\\pm0.9$\\\\ Proper motion DEC (mas/yr) & $-24.46\\pm0.8$\\\\ Spectral Type & F6IIIe \\\\ K (mag) & $4.98\\pm0.02$ \\\\ Distance (pc) & $145\\pm 15$ \\tablefootmark{a} \\\\ Age (Myr) & $2-20$ \\\\ Mass ($M_\\odot$) & $2.2\\pm0.3$ \\tablefootmark {a}\\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\tablefoot{ \\tablefoottext{a}{Based on new estimation of the stellar luminosity in Verhoeff et al. 2011.} } \\end{table} ",
        "conclusions": "Our observations using VLT/NaCo adaptive optics thermal and angular differential imaging reveal the disk surrounding HD\\,142527 at an unprecedented spatial resolution down to the subarcsecond level. The disk shows a vast and asymmetric cavity between $30~$AU and $145~$AU as well as two important dips towards North and South East. We also resolve a bright and inhomogeneous spiral arm towards the West direction, several fainter spiral arms and structured features rather than two elliptical facing arcs as proposed in previous studies. The observations in angular differential imaging allow us to derive unprecedented detection performance of point-sources down to planetary-mass at $40~$AU from the star. Monte-Carlo simulations combined with our detection map set stringent constraints on the presence of substellar companions orbiting HD\\,142527. We exclude the presence of brown dwarfs and massive giant planets in the outer part of the cavity and the presence of giant planets down to $5~$M$_{\\rm{Jup}}$ further out.  The observed structures indicate that the disk is in an active phase of global evolution. We have already discussed the possibility that the binary companion candidate might be highly responsible for the creation of these features. However, other processes can be invoked such as one or more unseen planets under formation in the disk.  Gravitational instability is known to create global spiral modes and in some cases fragmentation. However, given that the gas mass, as well as the surface density had been poorly constrained so far, it is not possible to know if the conditions are favorable for gravitational instability. Type II migration of a high mass planet can also be a mechanism able to produce such structures \\citep[e.g.][]{armitage05} in a gaseous disk. Thus, once again, constraints on the gas properties will allow to exclude or not the type II migration.  While this article was under review \\citet{casassus12} reported Gemini/NICI Ks and L$~\\!'$ observations of the outer part of the HD 142527 circumstellar environment. Given similar performance of NICI and NaCo, it gives us the opportunity to validate the observed features and excluded them to be instrumental or analysis artifacts. They confirmed the presence of several structures : the spiral arms, the asymmetric inner cavity, the two North and South South East dips and the inhomogeneity of the East side. However, we do not confirm in our images the presence of what they called a 'knot' along the North West arm. In their images, we also clearly see the offset of the main arm that we reported towards the West, even if they did not mentioned this peculiar morphology. Furthermore, they investigated the presence of a $10~$M$_J$ protoplanet at $90~$AU through hydrodynamic simulations to reproduce the morphology of the inner cavity and of the outer disk. Given our detection performance, such a companion would have likely been detected thus ruling out the input parameters of this carving planet. They argued themselves for additional ingredients within the simulations to reproduce all structures for which one single planet cannot explain. Nevertheless, their approach is very interesting and could be improved by taking into account our detection limits to set planet characteristics and the inner binary candidate companion within FARGO simulations.  From the present data, it is rather difficult to identify the cause of the non axisymmetric cavity and spiral arms due to the lack of knowledge on the gas properties. However, we consider that they may be due to at least one bound-embedded object. Further observations and important numerical investigations will be necessary to further test the origin of the highly structured disk of HD\\,142527."
    },
    "1208/1208.3141_arXiv.txt": {
        "abstract": "We present a new global model of the solar corona, including the low corona, the transition region and the top of chromosphere. The realistic 3D magnetic field is simulated using the data from the photospheric magnetic field measurements. The distinctive feature of the new model is incorporating the MHD Alfven wave turbulence. We assume this turbulence and its non-linear dissipation to be the only momentum and energy source for heating the coronal plasma and driving the solar wind. The difference between the turbulence dissipation efficiency in coronal holes and that in closed field regions is because the non-linear cascade rate degrades in strongly anisotropic (imbalanced) turbulence in coronal holes (no inward propagating wave), thus resulting in colder coronal holes with the bi-modal solar wind originating from them. The detailed presentation of the theoretical model is illustrated with the synthetic images for multi-wavelength EUV emission compared with the observations from SDO AIA and Stereo EUVI instruments for the Carrington rotation 2107. ",
        "introduction": "\\label{Intro} Results from Hinode observations have recently upped the estimates for the MHD wave energetics in the solar corona \\citep{dupo08}. Observed in the chromosphere, the magnetic field perturbations appear to be so powerful that even a $10\\div20\\%$ fraction of them, which propagate out from the Sun, carry a large enough energy to heat the solar corona and accelerate the solar wind. Even more promising, with the launch of the Solar Dynamics Observatory, we are now beginning to see observational hints of these ubiquitous waves in the transition region and low corona \\citep{mcintosh11}.  Even before these encouraging observations had been obtained, models which incorporated, or even were entirely based upon Alfv\\'en wave turbulence as the momentum and energy source were developed to describe the solar wind and coronal heating. Nowadays, {\\it developing the turbulence-driven global space weather model} becomes a problem tempting to be solved. \\subsection{Solar wind: can the turbulence-driven model compete with the semi-empirical one?} In \\cite{usma00}, a three-dimensional (3D) model for the solar wind was suggested in which the Alfv\\'en wave turbulence pressure served to accelerate the solar wind. The solar wind bi-modal structure as observed by Ulysses was successfully reproduced in the numerical simulation. However, the quantitative agreement of this model with long history of the solar wind observations at 1 AU is insufficient for global space weather simulations. The same criticism seems to be applicable to the more refined and physics-based  Alfv\\'en-wave-driven models of the solar wind \\cite{suzu05,verd10,osman11}.  Therefore, the semi-empirical approach so far is better suited for global space weather simulations. The most popular parameterization was adopted by \\cite{arge00} in their Wang-Sheeley-Arge (WSA) model, which well describes the solar wind parameters at 1~AU.  The semi-empirical ``synoptic'' formulae for the solar wind speed employ the solar magnetogram and the properties of the magnetic lines of the {\\em potential magnetic field} as recovered from the synoptic magnetogram data.  The empirical dependence of the solar wind properties from two input parameters: $\\theta$, which is the angular distance from the solar wind origin point till the nearest coronal hole boundary, and $f_{\\rm exp}$, which is the expansion factor, $f_{\\rm exp}=|{\\bf B}_{R = R_\\odot}|R_s^2/[|{\\bf B}_{R=2.5R_\\odot}|(2.5R_s)^2]$, is derived from the following two assumptions and  observations.  First, the slow solar wind is assumed to originate from the coronal hole boundary (small values of $\\theta$), whereas the fast solar wind originates from the central part of the coronal hole (large values of $\\theta$). Second, small coronal holes (having large values of $f_{\\rm exp}$) usually produce slower solar wind compared to large coronal holes (having small values of $f_{\\rm exp}$).  In \\cite{cohen07}, we coupled the semi-empirical WSA formulae to the global three-dimensional model for the solar corona and inner heliosphere within the Space Weather Modeling Framework (SWMF) \\cite[]{toth05}.  The WSA formulae were used as the boundary condition  for the model which had been coupled to the MHD simulator via the varied polytropic index distribution (see \\cite{roussev03b}). However, the physics of the Alfv\\'en wave turbulence has almost no intersection with the semi-empirical model of WSA.  {\\it In order to be competitive with semi-empirical model of the solar wind, the turbulence-driven model should quantitatively reproduce the solar wind variation from the coronal hole boundary to the coronal hole central part, similar to that parameterized by the expansion factor.} \\subsection{Coronal heating: can the turbulence-based model compete with {\\it ad hoc} heating functions?} Proceeding from the solar wind to the coronal heating, we can see again the disconnection between the physics-based models for the turbulent heating in the corona, one of the most advanced models of this kind being recently described in \\cite{cran10} (see also \\cite{tu97,hu00,Habb03,dmit02}) on one hand, and well established models with semi-empirical heating function, such as that presented in \\cite{lion01,rile06,tito08,lion09,downs10}, on the other hand. The heating function is applied to power the plasma in the solar corona with some heating rate, the distribution of which is given by some functions chosen not from some deep physical considerations, but chosen in an ad-hoc manner to provide a better agreement with observations.  These heating functions should be applied in conjunction with thermodynamic energy equation(s) in the Lower Corona (LC) model.  A direct accounting is needed of the transition region between the chromosphere and the corona, where non-ideal-MHD thermodynamic terms of energy transport, such as electron heat conduction, radiative losses, and coronal heating all become important. In \\citet{downs10} we used this thermodynamic MHD model to explore empirical parameterizations of coronal heating in the context of realistic 3D magnetic structures observed in the EUV and soft-X-Rays on Aug 27, 1996. Through direct comparison of synthetic observables to observations we demonstrated that this model can effectively capture the interplay between coronal heating and electron heat conduction.  The plasma parameters distribution obtained in this way may be successfully applied to generate synthetic EUV and X-ray images, which appear to be in a good agreement with those obtained with the EIT telescope onboard SOHO and SXT onboard Yohkoh (up to 2001). The observation synthesis capability has been extended to the major low corona imaging instruments available in space, namely STEREO/EUVI, SDO/AIA, and Hinode/XRT. The best agreement with observations is achieved with {\\it ad hoc} heating functions (such as the unsigned-flux-based heating model as discussed in Section 2.1).  {\\it In order to be competitive with ad hoc coronal heating models, the turbulence-driven model should quantitatively reproduce some successful heating function, that is the realistic wave dissipation should provide the same heating rate, as, for example, the unsigned-flux-based heating model \\cite{abbe07}.} \\subsection{Towards a Global Alfven-Wave-Turbulence-Driven MHD Model of the Solar Corona and Solar Wind} \\label{Instro_SC} In addition to the mentioned disconnection between physics-based and observations-driven models, there are two more contradictions to comment on. First, the quantitative Alfv\\'en-wave-driven models for the solar wind and for the coronal heating do not well conform with each other.  For example, the the coronal heating model of \\cite{cran10}, once applied to realistic 3D global model does not give realistic plasma parameters in the solar wind region (in the coronal holes).  Second, the increased estimates for the Alfv\\'en wave turbulence energetics in the Sun's proximity require us to revisit the models for the evolution of turbulence while the solar wind propagates towards 1~AU.  The revisited model should account for both {\\em the observed level and frequency spectrum of turbulence at 1~AU and the solar wind ion temperature resulting from the turbulence dissipation.}  With the advent of modern computational tools it is now becoming the norm to employ detailed 3D computer models as simulation tools that directly account for the inhomogeneous nature of the Sun-Heliosphere environment. The key advantage of this approach is the ability to compare and validate model results through direct comparisons to all kinds of the  observational data listed above: for the solar wind and turbulence characteristics at 1 AU (see also \\cite{jin12}) and for EUV and X-ray images of the solar corona. \\subsection{Goal and Content of the Paper} In the current paper, we get rid of the {\\it ad hoc} heating functions and parameterize the coronal heating in the LC in terms of the Alfven wave turbulence dissipation.  The theoretical model for this approach is summarized in Section 2. In subsections 2.1, 2.3 we demonstrate the possibility to parameterize a popular and successful semi-empirical heating model based on unsigned magnetic flux in terms of the Alfv\\'en wave turbulence dissipation. While choosing the way to parameterize heating, one may try to vary both the boundary condition for the wave energy flux (the Poynting flux) inflowing to the solar corona from the solar surface and the dissipation length for turbulence. In subsection 2.2 we show that the boundary condition is strongly restricted which reduces the model uncertainty. With the pre-specified boundary condition, the drastic difference in the plasma heating mechanism between open and closed field regions should be caused by a difference in the wave dissipation efficiency. In subsection 2.4 we suggest and describe the physical mechanism of a nonlinear interaction between {\\it oppositely propagating waves}. The degraded intensity of the inward propagating waves may be responsible for the reduction in the turbulence dissipation rate in the coronal holes thus resulting in the bimodal solar wind structure.  In Section 3 we summarize a computational model and the code we use to simulate the state of the solar corona. Section 4 presents the simulation results for CR2107 and their comparison with EUV images. In the Conclusion we discuss the plans for a future. ",
        "conclusions": ""
    },
    "1208/1208.5417_arXiv.txt": {
        "abstract": "{ Atmospheric conditions at the site of a cosmic ray observatory must be known well for reconstructing observed extensive air showers, especially when measured using the fluorescence technique. For the Pierre Auger Observatory, a sophisticated network of atmospheric monitoring devices has been conceived. Part of this monitoring was a weather balloon program to measure atmospheric state variables above the Observatory. To use the data in reconstructions of air showers, monthly models have been constructed. Scheduled balloon launches were abandoned and replaced with launches triggered by high-energetic air showers as part of a rapid monitoring system. Currently, the balloon launch program is halted and atmospheric data from numerical weather prediction models are used. A description of the balloon measurements, the monthly models as well as the data from the numerical weather prediction are presented. \\PACS{ {95.85.Ry}{Cosmic rays --} {95.55.Vj}{ Cosmic ray detectors --} {92.60.-e}{ Properties and dynamics of the atmosphere} } } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.0004_arXiv.txt": {
        "abstract": "\\mysou is an obscured accreting X-ray pulsar discovered by \\integral on 2012 February 29. We report on our 20\\,ksec \\chandra-High Energy Transmission Gratings Spectrometer observation of the source performed on 2012 March 17, on two short contemporaneous \\swift observations, and on our two near-infrared ($\\mathrm{K}_s$, $\\mathrm{H}_n$, and $\\mathrm{J}_n$) observations performed on 2012 March 13 and March 26. We determine the most accurate X-ray position of \\mysou, $\\alpha_\\mathrm{J2000}=18^\\mathrm{h}17^\\mathrm{m}52^\\mathrm{s}.18$, \\mbox{$\\delta_\\mathrm{J2000}=-16^\\circ 21\\darcmin\\, 31\\darcsec.68$} (90\\% uncertainty of 0\\darcsec.6). A strong periodic variability at 11.82\\,s is clearly detected in the \\chandra data, confirming the pulsating nature of the source, with the lightcurve softening at the pulse peak.  The quasi-simultaneous \\chandra-\\swift spectra of \\mysou can be well fit by a heavily absorbed hard power-law ($N_\\mathrm{H} =2.2 \\pm0.3 \\times 10^{23}\\,\\mathrm{cm}^{-2}$, and photon index $\\Gamma = 0.4\\pm0.1$) with an average absorbed 2--8\\,keV flux of $1.4\\times 10^{-11}\\,\\mathrm{erg}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}$. At the \\chandra-based position, a source is detected in our near infrared (NIR) maps with $\\mathrm{K}_s = 13.14 \\pm 0.04$\\,mag, $\\mathrm{H}_n = 16\\pm0.1$\\,mag, and no $\\mathrm{J}_n$ band counterpart down to $\\sim$18\\,mag. The NIR source, compatible with 2MASS~J18175218$-$1621316, shows no variability between 2012 March 13 and March 26. Searches of the UKIDSS database show similar NIR flux levels at epochs six months prior to and after a 2007 February 11 archival \\chandra observation where the source's X-ray flux was at least 87 times fainter. In many ways \\mysou is unusual: its combination of a several week long outburst (without evidence of repeated outbursts in the historical record), high absorption column (a large fraction of which is likely local to the system), and 11.82\\,s period does not fit neatly into existing X-ray binary categories. ",
        "introduction": "\\label{sec:intro} \\setcounter{footnote}{0}  One of the great contributions of the \\integral satellite has been its scans of the Galactic plane and bulge, which have led to the discovery of a number of previously unknown transient X-ray binary sources \\citep{bodaghee07}. Many of these sources have high columns of absorbing material along the line of sight, making the \\integral hard X-ray response crucial for their discovery. In numerous cases a significant fraction of these columns are \\emph{intrinsic} to the systems in question. In fact, \\integral has suggested possible interesting relationships between absorption and binary properties: a neutral column/orbital period anti-correlation (not unexpected as short periods increases the likelihood that an X-ray source is more deeply embedded within the wind/outflow from the secondary), and a neutral column/X-ray pulsar spin period correlation \\citep[which, if real, lacks an explanation;][]{bodaghee07}.  X-ray binaries themselves represent a wide variety of types of systems: High Mass X-ray Binaries (HMXB), which may be either wind fed or Roche Lobe Overflow fed, Be/X-ray binaries, and accreting msec X-ray pulsars \\citep[e.g.,][]{paizis05b}, among other types of systems. Outbursts can be very rapid, lasting only a few days in the case of the so-called Supergiant Fast X-ray Transients (SFXT; \\citealt{negueruela06,sidoli11}, and references therein), or last several weeks in the cases of Be/X-ray binaries and accreting msec X-ray pulsar systems. \\integral, in combination with observations at other wavelengths, has been crucial in identifying the natures of these sources.  Since 2005 we have had a \\chandra-followup program to observe previously unknown X-ray sources discovered by \\integral, in order to aid in localizing and identifying these sources for multi-wavelength follow-up observations. On 2012 February 29 (MJD 55986) \\integral discovered a new hard X-ray transient \\mysou \\citep{tuerler12}. The best source position was determined with the two JEM--X instruments with an associated uncertainty of \\mbox{1\\darcmin.5}. The combined JEM--X and IBIS/ISGRI spectrum ($F_\\mathrm{3-50\\,keV}=1.0\\times 10^{-9}\\,\\mathrm{erg}\\,\\mathrm{cm}^{-2}\\,\\mathrm{s}^{-1}$) could be described by a cut-off power-law ($\\Gamma=-0.5\\pm0.5$, $E_\\mathrm{cutoff}= 4.9^{+1.5}_{-0.9}\\,\\mathrm{keV}$, where photon flux per unit energy is $\\propto E^{-\\Gamma}\\exp(-E/E_\\mathrm{cutoff})$) plus a broad Gaussian absorption line ($E_\\mathrm{c}={20.8}^{+1.4}_{-1.8}$\\,keV).  According to the authors, if this line is interpreted as a cyclotron resonant scattering feature, \\mysou could be a High Mass X-ray Binary (HMXB) pulsar with a magnetic field of about $1.7 \\times 10^{12}$\\,Gauss. The pulsating nature was confirmed by \\swift and \\textit{Fermi}/GBM detections of a strong signal at $P=11.82\\pm0.01$\\,s \\citep{halpern12,li12,finger12}. The \\swift observations performed on February 29 confirmed the hard spectrum of \\mysou ($\\Gamma=0.4\\pm0.4$) as well as a high absorbing column density, $N_\\mathrm{H}=(11.0\\pm2.0)\\times10^{22}\\,\\mathrm{cm}^{-2}$ \\citep{li12}.  A refined position with respect to the \\integral discovery ($1\\darcmin.5$ uncertainty) was reported by \\cite{li12} using \\swift data. That position ($2\\darcsec.2$ uncertainty) was consistent with the position measured by \\integral and was compatible with an infrared candidate counterpart, 2MASS~J18175218$-$1621316. However, this was not confirmed by \\cite{halpern12} who reported a \\swift-based source position 4\\darcsec.4 away (no error given), hence the proposed association with the 2MASS source needed confirmation.  The day after the discovery of \\mysou, we triggered our approved \\chandra target of opportunity program.  The observation was originally set to occur on 2012 March 11, but due to a strong solar flare it was postponed until 2012 March 17.  Thanks to the excellent \\chandra imaging and astrometry, an X-ray position with a \\mbox{0\\darcsec.6} (90\\% confidence level) uncertainty was reported \\citep{paizis12}. This new \\chandra-based position was 0\\darcsec.37 away from that obtained by \\citet{li12} and 0\\darcsec.09 from the 2MASS~J18175218$-$1621316 source ($\\mathrm{K}_s$=13.14\\,mag). The { preliminary} \\chandra position confirmed this latter source as the best candidate counterpart to \\mysou. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{\\mysou: X-ray/NIR}\\label{sec:nature}  2MASS~J18175218$-$1621316 (inner red circle in Fig.~\\ref{fig:nir}) is the only source from the 2MASS All-Sky Point Source Catalog that is within 1\\arcsec of the \\chandra-determined position of \\mysou, and the 2MASS point source density within 60\\,$\\arcsec$ of this position is $0.01/$arcsec$^2$. Given that the \\chandra error circle for \\mysou is $\\approx 1$\\,arcsec$^2$, the simplest hypothesis is that this 2MASS source (with a $\\sim 0.1\\,\\arcsec$ error radius) is the actual NIR counterpart to the X-ray source.  There still remains, however, an $\\aproxgt 1\\%$ probability that the 2MASS source is a chance coincidence.  As we discuss below in \\S\\ref{sec:hmxbs}, if \\mysou lies between 1--10\\,kpc distant, a wide range of companion types plausibly would have NIR flux levels and colors of this 2MASS source, even in the presence of large extinction.  In what follows, we therefore adopt the hypothesis that 2MASS~J18175218$-$1621316 is indeed the NIR counterpart to \\mysou.  The 2MASS catalogue provides a source $\\mathrm{K}$ magnitude of $13.14\\pm0.04$\\,mag, to be compared to our constant NIR observed $13.14 \\pm 0.04$\\,mag and the UKIDSS $\\mathrm{K}$ magnitude of 13.007-13.043\\,mag. Hence, though the X-ray outburst decay inferred from \\swift monitoring ($\\sim$22\\,days, \\citealt{bozzo:12a,li:12a}) may resemble a low mass X-ray binary outburst type of event (accretion disk instability), the essentially constant level of the NIR emission (in the 2MASS and UKIDSS catalogues and in our observations during the outburst) suggests that we are seeing the stable high mass companion at $\\mathrm{K}_s \\sim 13$\\,mag.  The $\\mathrm{J}$-band upper limits from our observations and UKIDSS further indicate that this companion is highly obscured.  The heavily absorbed ($N_\\mathrm{H} =(2.2 \\pm 0.3) \\times 10^{23}\\,\\mathrm{cm}^{-2}$) and hard spectrum ($\\Gamma = 0.4\\pm0.1$) obtained in the joint \\chandra-\\swift spectral fit (Section~\\ref{sec:swiftchandraspe}), together with the 11.82\\,s pulsation (Section~\\ref{sec:chandravar}) and the putative cyclotron line discussed by \\citet{bozzo:12a} and \\citet{li:12a} point to an identification of \\mysou as an HMXB hosting a neutron star (NS).  The expected neutral column along this line of sight is $\\approx 10^{22}\\,\\mathrm{cm}^{-2}$ \\citep{dickey90,kalberla:05a}, which is more than a factor of ten smaller than the measured column. A substantial fraction of the measured neutral column is therefore likely local to the system and associated with an atmosphere/wind/accreting matter coming from the secondary. Associated with this local absorption we expect an Fe\\,K$\\alpha$ emission line with EW$\\approx 3.3\\,( {\\rm N_H}/10^{22}\\,{\\rm cm^{-2}})$\\,eV $\\approx 73$\\,eV, as has been found empirically in \\chandra studies of HMXB \\citep{torrejon10a}. Such a Fe line strength is also expected under the approximation that the obscuring material is spherically distributed about the X-ray source \\citep{kallman04a}. Monte Carlo simulations of a power-law source embedded within an absorbing cloud show the line equivalent width to be between 50\\,eV and 100\\,eV for material of solar and twice solar abundances (Eikmann, priv.\\ comm.). For these reasonable assumptions, the lack of a detectable Fe K$\\alpha$ line is therefore consistent with the expectations and it can be assumed that the bulk of the observed neutral column is local to the system.  Figure~\\ref{fig:timing} (middle and right panels) shows that the softer 1--5\\,keV folded lightcurve has greater modulation than the harder 5--8\\,keV lightcurve. That is, the lightcurve softens at the peak of the 11.82\\,s pulse.  Such an energy dependent pulse profile is seen in many magnetized neutron stars, where the stronger variability at lower energies is typically attributed to absorption in the accretion column, while the (typically sinusoidal) emission pattern at harder energies is due to emission from the bottom of the accretion column \\citep[][and references therein]{caballero:11a}.   \\subsection{The Nature of \\mysou: an HMXB?}\\label{sec:hmxbs}  HMXBs are known to be divided into different groups according to the nature of the companion, formation and geometry of the system \\citep[e.g.,][and references therein]{chaty11}. They include BeHMXBs X-ray transients, usually hosting a NS on a wide eccentric orbit around a B0--B2e spectral type donor star with a decretion disk; and supergiant HMXBs (sgHMXBs), hosting a NS on a circular orbit around a supergiant OB star. sgHMXBs can be either Roche lobe overflow sources (RLO; rare transient sources with outburst luminosities of $L_\\mathrm{X}\\sim10^{38}\\,\\mathrm{erg}\\,\\mathrm{s}^{-1}$), or wind-fed systems (e.g., the majority of sgHMXB, with persistent luminosities of $L_\\mathrm{X}\\sim10^{35-36}\\,\\mathrm{erg}\\,\\mathrm{s}^{-1}$).  \\begin{figure}\\centering \\includegraphics[width=0.4\\textwidth]{corbet4.eps} \\caption{Corbet diagram (NS $P_\\mathrm{spin}$ versus system $P_\\mathrm{orb}$; \\citealt{corbet86}) showing the three HMXB populations: BeHMXB, sgHMXB-Roche lobe overflow and sgHMXB-wind fed. See text. \\mysou ($P_\\mathrm{orb}$ unknown) lies along the horizontal blue line at $P_\\mathrm{spin}\\sim$12\\,s.  Adapted from a figure by J. A. Zurita Heras, as originally presented by \\protect{\\citet{chaty11}}.} \\label{fig:corbet} \\end{figure}  An efficient way to visualize all such sources is via the so-called Corbet diagram \\citep{corbet86}, where the NS spin period, $P_\\mathrm{spin}$, is plotted versus the system orbital period, $P_\\mathrm{orb}$ (Fig.~\\ref{fig:corbet}). In general, in such a diagram BeHMXBs populate the region with higher $P_\\mathrm{orb}$ (roughly $P_\\mathrm{orb}> 20$\\,days), while sgHMXBs populate the lower part ($P_\\mathrm{orb} < 20$\\,days). $P_\\mathrm{spin}$ ranges from about 1\\,s to $10^{3}$\\,s for both classes of sources (Fig.~\\ref{fig:corbet}; see also \\citealt{chaty11, sidoli11}). The majority of transient sources with spin periods $\\aproxlt 20$\\,s are in fact BeHMXB.  In the last decade, two new classes of wind accreting pulsars in sgHMXBs have been discovered \\citep{bodaghee07}: obscured, fairly persistent sources with a huge intrinsic local extinction, and sgHMXBs exhibiting fast and transient outbursts, known as Supergiant Fast X-ray Transients \\citep[SFXT, see e.g.,][for a recent review, especially Fig.~2 therein]{sidoli11}. The neutral column of \\mysou is typical of this latter category of sources, but its spin period -- and prolonged outburst -- is not.  Based on the above discussion, \\mysou does not clearly fall into any of the above classes. Its transient nature seems to exclude the highly obscured and standard wind-fed systems that are fairly persistent. A few week long transient outburst of \\mysou could suggest that it is a BeHMXB, with its 11.82\\,s spin period placing it at the lower end of orbital periods in the Corbet diagram, i.e., several tens of days (20--80\\,days).  Although many Be X-ray binaries show periodic outbursts, many other Be show long periods of quiescence. A good example is A0535+26, which has had only five major outburst phases since its 1975 discovery, which were separated by many years of quiescence \\citep[e.g., between 1994 February and 2005 May, see][and references therein]{caballero:07a}.  On the other hand, despite its transient nature \\mysou would be unusual for a SFXT, since these sources are known to have very fast outbursts, on the order of hours or a few days at most, as seen by \\integral/IBIS. Fig.~\\ref{fig:lcr_all} shows that \\mysou is detected for at least three revolutions, i.e., about 12 days. An extremely nearby SFXT, however, may not be excluded.  RLO sources, i.e., transient sgHMXBs strictly filling their Roche-Lobe (pure accretion disk accretors) are rare, since the mass transfer is highly unstable and the accretion should only last for a few thousand years. Instead, there are HMXBs exhibiting beginning atmospheric Roche Lobe overflow, where the massive star does not fill its Roche Lobe, but the stellar wind follows the Lagrange equipotentials, and accumulates, forming an accretion disk \\citep{bhattacharya1991}. This situation is more stable, but still rare due to the required configuration of stellar radius, orbital distance and mass ratio. We know only three such systems in total, hosting NSs: LMC~X-4, Cen~X-3 and SMC~X-1. \\mysou could be the fourth source belonging to this class.  \\begin{figure}\\centering \\includegraphics[width=0.36\\textwidth,angle=90]{chatyII.ps} \\caption{Plot of expected H$-$K colors vs.\\ absolute K magnitude for different stellar types. Superimposed on this are points (asterisks) for the NIR measurements of \\mysou, assuming three different visual extinctions ($A_\\mathrm{V}=$47, 50, and 53\\,mag) for assumed distances ranging from 1--10\\,kpc.} \\label{fig:distance} \\end{figure}  Finally, we consider the possible stellar companions in \\mysou, based upon the NIR measurements coupled with assumptions about possible distances and assumed extinctions. In Fig.~\\ref{fig:distance}, we show absolute K-magnitude vs.\\ H$-$K colors, for a series of stellar sequences. Overplotted on this are the absolute K-magnitude and H$-$K colors for the counterpart of \\mysou given our measurements, and assuming distances ranging from 1--10\\,kpc.  We estimate the extinction towards the source from the relationship of \\citet{predehl95}, modified to account for the fact that absorption model of \\citet{wilms00} fits neutral columns $\\approx 30\\%$ larger than the model used by \\citet{predehl95}. Specifically, we assume $A_\\mathrm{V} \\sim N_\\mathrm{H}/2.7\\times10^{21}\\,\\mathrm{cm}^{-2}$. The 21\\,cm column translates to an extinction of $A_\\mathrm{V}=3.7$\\,mag, while the full fitted neutral column yields $A_\\mathrm{V}\\sim 82$\\,mag. Note, however, that the \\citet{predehl95} relationship is based on measurements of X-ray dust scattering halos. Source intrinsic absorption in the stellar wind would happen in a medium that does not contain dust and thus for such a case the \\citet{predehl95} relationship overestimates the $A_\\mathrm{V}$.  Figure~\\ref{fig:distance} shows that the range of extinctions from $A_\\mathrm{V} =47$--52\\,mag is the only range where our NIR measurements intersect standard stellar sequences. While \\mysou could be consistent with a Low Mass X-ray Binary (LMXB), with a red giant companion ranging from K-type ($A_\\mathrm{V} \\sim 49$, at 1\\,kpc) to M-type ($A_\\mathrm{V} \\sim 47$, at 10\\,kpc), a LMXB nature seems unlikely, given the 11.82 s pulsation and very hard X-ray spectrum of \\mysou. More realistic models for \\mysou are that it is a HMXB. Supergiant companions can range from a close by OB-type ($A_\\mathrm{V} \\sim 52$, at 2\\,kpc) to a further away A-type ($A_\\mathrm{V} \\sim 49$, at 7\\,kpc). Even a nearby B-type main sequence star is allowed ($A_\\mathrm{V} \\sim 52$), at 1\\,kpc.  In all these cases, $A_\\mathrm{V}$ is $\\sim$50\\% of what is obtained by calculating $A_\\mathrm{V}$ from the X-ray $N_\\mathrm{H}$. This result argues that about half of the observed column is caused by the neutron star being embedded in the stellar wind of an early-type donor, while the other half of the column would be in the (dusty) interstellar medium close to the binary. The deduced column for the wind absorption, $N_\\mathrm{H, wind}\\sim 10^{23}\\,\\mathrm{cm}^{-2}$, is consistent with that seen in a number of HMXB.  For example, the XRB GX 301$-$2, which orbits a B1 Ia hypergiant in a 41.5\\,d ellipitical orbit, has an $N_\\mathrm{H}$ that is strongly variable, with $N_\\mathrm{H}\\sim 10^{24}\\,\\mathrm{cm}^{-2}$ during the pre-periastron flare \\citep{fuerst:11a} and $N_\\mathrm{H}\\sim 10^{23}\\,\\mathrm{cm}^{-2}$ and lower after the periastron passage \\citep{suchy:12a}.  Unfortunately, we lack detailed observations over the course of the outburst that might have revealed variations of the neutral column.  The nature of \\mysou at this time still remains ambiguous."
    },
    "1208/1208.5551_arXiv.txt": {
        "abstract": "The internal structures and compositions of Uranus and Neptune are not well constrained due to the uncertainty in rotation period and flattening, as well as the relatively large error bars on the gravitational coefficients. While Uranus and Neptune are similar in mass and radius, they differ in other physical properties such as thermal emission, obliquity, and inferred atmospheric enrichment. In this letter we consider the uncertainty in the planetary rotation periods, show that rotation periods more consistent with the measured oblateness imply that Uranus and Neptune have different internal structures, and speculate on the source of that difference.  We conclude that Uranus and Neptune might have very different structures and/or compositions despite their similar masses and radii.  We point out that understanding these differences can have important implications for our view of the formation and evolution of Uranus and Neptune as well as intermediate-mass extra-solar planets in general. ",
        "introduction": "Uranus and Neptune are often thought of as twin planets. They both formed in the outer solar system, and their masses and radii are very similar, as are their gravitational moments.  Even their rotation periods are within 10\\% of each other. Yet, there are important differences between these planets. Uranus' mass is slightly smaller than Neptune's, but its radius is somewhat larger, so the difference in mean density, is considerable. Indeed interior models of Uranus and Neptune (Podolak et al. 1995; Helled et al. 2011) indicate that Neptune's outermost envelope is more enriched in high-Z material than that of Uranus.  In addition, Neptune has an internal heat source, while Uranus is in equilibrium with solar insolation (Pearl et al. 1990; Pearl \\& Conrath 1991) suggesting that Uranus' interior may not be fully convective, and/or that it contains compositional gradients which hinder convection. The difference in thermal flux is indicated from calculations of Uranus' thermal evolution which find that unlike Neptune, Uranus cannot reach its measured intrinsic luminosity within solar system's age if an adiabatic interior is assumed (Fortney \\& Nettelmann 2010, Fortney et al., 2011). Another distinct feature of Uranus is its large axial tilt. This was almost certainly caused by a dramatic event in the early history of the solar system (Safronov 1966), although an alternative explanation for Uranus' tilt is orbital migration ({Bou\\'{e}} \\& {Laskar}, 2010).  Stevenson (1986) suggested that this same dramatic event might have caused the necessary compositional gradients for inhibiting convection. We return to this point later.  Despite the available measurements, the internal structures of Uranus and Neptune are not well constrained, and as a result, their bulk compositions are essentially unknown (Podolak et al. 1995; Marley et al. 1995; Helled et al. 2011). In fact, it is still unclear what the mass fraction of water is in Uranus and Neptune, despite their categorization as `icy planets'. \\par  Recently, Nettelmann et al. (2012), showed that by assuming modified solid-body rotation periods for Uranus and Neptune, the derived internal structures can change considerably. It was concluded  that an uncertainty in the planetary rotation period can be crucial for constraining the internal structure. In this letter, we investigate further how the uncertainties in the gravity field, shape, and rotation period might affect the inferred planetary interior. ",
        "conclusions": "The three measured parameters that are used to determine the structure of Uranus and Neptune's ($J_2, q, \\omega$) are not self-consistent and any chosen pair predicts a different inertia factor for the planets.  The choice of $q$ and $J_2$ which are the ones used in interior modeling, gives very similar inertia factors for Uranus and Neptune implying that their internal structures are alike. However, when we use the HAS10  rotation periods, the parameters $q$, $f$ and $J_2$ are more self-consistent and Uranus is found to be more centrally condensed than Neptune. As a result, Uranus and Neptune may differ substantially in their internal structures (Nettelmann et al., 2012).  Two important results derive from our study.  First, one must be cautious in drawing conclusions about planetary structure simply on the basis of a mass classification.  Uranus and Neptune differ by only 10\\% in mass and only 3\\% in radius, but they may have very different interiors.  Second, these differences may be caused by secondary effects, such as the impact parameter of a giant impact.  These effects have important consequences for the classification of intermediate-mass bodies, and more detailed studies are necessary in order to evaluate the differences we might expect in the structure and composition of the many intermediate-mass exoplanets."
    },
    "1208/1208.5598_arXiv.txt": {
        "abstract": "{Measuring time delays between the multiple images of gravitationally lensed quasars is now recognized as a competitive way to constrain the cosmological parameters, and it is complementary with other cosmological probes. This requires long and well sampled optical light curves of numerous lensed quasars, such as those obtained by the COSMOGRAIL collaboration. High-quality data from our monitoring campaign call for novel numerical techniques to robustly measure the delays, as well as the associated random and systematic uncertainties, even in the presence of microlensing variations. We propose three different point estimators to measure time delays, which are explicitly designed to handle light curves with extrinsic variability. These methods share a common formalism, which enables them to process data from $n$-image lenses. Since the estimators rely on significantly contrasting ideas, we expect them to be sensitive to different bias sources. For each method and data set, we empirically estimate both the precision and accuracy (bias) of the time delay measurement using simulated light curves with known time delays that closely mimic the observations. Finally, we test the self-consistency of our approach, and we demonstrate that our bias estimation is serviceable. These new methods, including the empirical uncertainty estimator, will represent the standard benchmark for analyzing the COSMOGRAIL light curves.} ",
        "introduction": "In the era of precision cosmology, in which a concordance model seems to fit independent observations, it is of utmost importance to both compare and combine all possible methods that constrain cosmological parameters. Comparing them yields an invaluable cross-check of the methods and the model. Combining them allows breaking the degeneracies inherent in single techniques. Probes including baryonic acoustic oscillations, weak lensing, supernovae, and cosmic microwave background measurements fit in this context exactly.  Also among these probes is the so-called ``time-delay method'', first proposed by \\citet{Refsdal:1964vh} to measure cosmological distances independently of any standard candle. In practice, the method uses strongly lensed quasars with significant photometric variability. Photons emitted by the source quasar propagate towards us along different optical paths, resulting in multiple images. The light travel times associated to these images differ due to (i) the different path lengths, and (ii) the different Shapiro delays induced by the gravitational field of the lensing galaxy. As a consequence, the same quasar variability is seen with distinct time shifts in the light curves of the multiple quasar images. This paper presents methods of inferring the relative time delays between the quasar images, from such resolved, i.e. unblended, light curves.  Measured time delays, in combination with deep HST imaging and dynamical information on the lensing galaxy lead to competitive measurement of the Hubble constant \\ho \\citep[e.g.,][]{Suyu:2009ig, Suyu:2010fq}. The complementarity between quasar time delays and several other cosmological probes has been illustrated recently by \\citet{Linder:2011cs}, who points out that the dark-energy figure of merit of a combination of Stage III experiments is improved by a factor of 5 if 150 quasar time delays are added. This also holds if the Universe is not assumed to be flat. It is noteworthy that adding this time delay information is very cheap compared to other Stage III or IV projects.  The COSMOGRAIL collaboration has now gathered almost a decade of photometric points for about 30 lensed quasars. With such data, the time delays can in most cases be seen clearly ``by eye''. The data analysis is no longer about sorting out which time delay is the best among several plausible yet incompatible possibilities, but rather about performing an accurate measurement of the delay that can be reliably used for cosmology. New \\emph{curve-shifting techniques} must be devised to extract the delays from such curves, which sometimes include a thousand points and typically display substantial microlensing variability due to stars in the lensing galaxy.  In this paper we present three independent curve-shifting algorithms that can deal with extrinsic variability. Our motivation behind the development of \\emph{several} techniques is to provide a range of methods that rely on different principles. While the methods might not be free of systematics, we expect them to be biased in different ways, and we devote a large part of this work to estimating comprehensive error bars. Comparing the results from different curve-shifting techniques will allow us, in particular, to systematically cross-check our quantification of the biases.  Our paper is structured as follows: Section \\ref{curves} gives an overview of features to be expected in light curves, most of them complicating the time delay extraction problem. We then present the point estimation formalism that is common to our curve-shifting techniques in Section \\ref{techniques}. Sections \\ref{splines} to \\ref{disp} describe the three techniques, and we explain how we consistently compute error bars for each time delay and technique in Section \\ref{errorbars}. We compare our techniques and the associated uncertainty estimates in Section \\ref{test}, using a set of simulated light curves with known time delays. Finally, we present a summary and our conclusions in Section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this paper, we describe three independent ``curve-shifting techniques'' to measure time delays between resolved light curves of gravitationally lensed quasars. All these methods address the presence of variable microlensing in the light curves and can be applied to lens systems with any number of images. \\begin{enumerate} \\item {\\bf The free-knot spline technique} simultaneously fits one common \\emph{intrinsic} spline and independent smoother \\emph{extrinsic} splines to the light curves. The curves are shifted in time so as to optimize this fit. \\item {\\bf The regression difference technique} shifts regressions of the curves to minimize the \\emph{variability} of the differences between them. It is nearly parameter free and does not require an explicit model for the microlensing variability. \\item {\\bf The dispersion-like technique} shifts the curves so as to minimize a measure of the \\emph{dispersion} between the overlapping data points. This method has no explicit model for the common intrinsic variability of the quasar, but it involves polynomial models for the extrinsic variability. It has previously been applied in \\citet{Courbin:2011bl}. \\end{enumerate} A common point of the methods is that they yield point estimates (i.e., single values) for the time delays in a self-consistent way by sharing the formalism of time \\emph{shifts} described in Section \\ref{techniques}.  In addition, we present a Monte Carlo approach to estimate the uncertainty of each time delay measurement, including both random and systematic errors. This procedure is based on synthetic curves that try to mimic as much information about the intrinsic and extrinsic variability as the observations unmistakenly reveal. Provided that we accept the generative model of the synthetic curves, the curve-shifting techniques themselves are reduced to ``recipes''. Given a set of light curves, we can select methods based solely upon their empirical performance. This effectively shifts the requirement of formal justification from the curve-shifting techniques to the synthetic curves on which these techniques are evaluated. As a consequence, the techniques can even be fine tuned for each data set.  Finally, we verified the self-consistency of our time delay uncertainty estimation using a trial set of artificial light curves. The availability of three different curve-shifting techniques allows consistency checks of our \\emph{bias} determination to be performed when analyzing real observations, i.e., data without known time delays. In other words, we acknowledge that any curve-shifting technique may display residual biases due, for example, to particular patterns of slow microlensing variability, but we provide a means to evaluate this bias.  The methods described in this paper will be used as a standard benchmark to obtain time delay estimations from the light curves of the COSMOGRAIL monitoring program. We implemented the curve-shifting techniques, the error bar estimation, and the generation of all the figures of this paper in the form of a modular \\verb+python+ toolbox. This package, called PyCS, as well as a tutorial including the trial curves of Section \\ref{test} are available from the COSMOGRAIL website\\footnote{\\url{http://www.cosmograil.org}}."
    },
    "1208/1208.5284_arXiv.txt": {
        "abstract": "{} {We study the production of VHE emission in blazars as a superposition of a steady component from a baryonic jet and a time-dependent contribution from an inner $e^-e^+$ beam launched by the black hole.} {Both primary relativistic electrons and protons are injected in the jet, and the particle distributions along it are found by solving a one-dimensional transport equation that accounts for convection and cooling. The short-timescale variability of the emission is explained by local pair injections in turbulent regions of the inner beam.} {For illustration, we apply the model to the case of PKS 2155-304, reproducing a quiescent state of emission with inverse Compton and synchrotron radiation from primary electrons, as well as proton-proton interactions in the jet. The latter also yield an accompanying neutrino flux that could be observed with a new generation km-scale detector in the northern hemisphere such as KM3NeT.} {} ",
        "introduction": "Blazars are the AGNs in which the jet points mainly in the direction of the line of sight. They exhibit the most extreme high-energy phenomena of all AGNs. Their spectral energy distributions (SEDs) are characterized by nonthermal continuum spectra with a broad low-frequency component from X-rays to $\\gamma$-rays. Blazars show rapid variability across the entire electromagnetic spectrum. Variability at high energies on timescales of a few minutes has been observed for some of them, such as PKS 2155-304 (e.g. Aharonian et al. 2006).  In this work we present a two-{component} jet model with both relativistic leptons and hadrons to explain the high-energy emission from these objects. The basic scenario consists of a steady baryonic jet launched by the accretion disk, and an $e^+e^-$ beam launched by the black hole ergosphere. The quiescent component of the signal is assumed to be produced by the jet, while the variable component is due to shocked regions in the inner $e^+e^-$ beam. Inhomogeneities and turbulence can be generated by Kelvin-Hemholtz instabilities. In Sect. 2 we describe the basics of the model. Its application to PKS 2155-304 is presented in Sects. 3 and 4 for the quiescent and variable emission, respectively. In Sect. 5 we focus on the neutrino output expected for the same blazar, analyzing the detectability with a next generation neutrino telescope such as KM3NeT. We finish in Sect. 6 with a discussion. ",
        "conclusions": "We have implemented a two-{component} model to study the emission of the blazar PKS 2155-304. The quiescent state of electromagnetic emission was associated with the contribution produced in a heavy barionic jet, in which both electrons and protons can be accelerated. Specifically, this low state of emission is dominated by synchrotron and inverse Compton interactions of the primary electrons, and also by $pp$ interactions in the jet. The time-dependent contribution {is produced by multiple shocks in} an internal electron-positron beam, surrounded by the jet. We showed that if shocks are injected in this beam, as would be expected as a consequence of Kelvin-Hemlholtz instabilities, then a variable emission can be generated via SSC interactions, giving rise to gamma-ray light curves similar to those observed by HESS.  Other important output predicted by the present model is in the form of VHE neutrinos produced by $pp$ interactions. If this accompanying flux remains constantly produced over a period of four years, the detection of at least one event from PKS 2155-304 should be guaranteed in a detector like KM3NeT with a probability very close to one. This would be a very important piece of evidence for hadronic acceleration in the source. Conversly, if the neutrinos are not detected at this level, this would be evidence of a leptonic-dominated source."
    },
    "1208/1208.1584_arXiv.txt": {
        "abstract": "We analyze the optical spectra of massive (log $M_*/M_\\odot > 11.4$) radio-loud galaxies at $z \\sim$ 0.2 and $z \\sim$ 0.6. Our samples are generated by crossmatching the SDSS DR7 and BOSS spectroscopic galaxy catalogues with the FIRST and NVSS radio continuum surveys. By comparing stellar population parameters of these radio-loud samples with radio-quiet control samples matched in stellar mass, velocity dispersion and redshift, we investigate how the presence of a radio-emitting jet relates to the recent star formation history of the host galaxy. We also investigate how the emission-line properties of the radio galaxies evolve with redshift by stacking their spectra. Our main results are the following.  (1) Both at low and at high redshift, half as many radio-loud as radio-quiet galaxies have experienced significant star formation in the past Gyr. This difference in star formation history is independent of the luminosity of the radio AGN, except at radio luminosities greater than $10^{25.5}{\\rm W~Hz}^{-1}$, where it disappears. (2) The Balmer absorption line properties of massive galaxies that have experienced recent star formation show that star formation occurred as a burst in many of these systems. (3) Both the radio and the emission-line luminosity of radio AGN evolve significantly with redshift. The average \\oiii\\ rest equivalent width increases by 1 dex from $z$ = 0.2 to $z$ = 0.6, and emission line ratios change from LINER-like at low redshift to Seyfert-like at high redshift. However, radio galaxies with similar stellar population parameters, have similar emission-line properties both at high- and at low-redshift.  These results suggest that massive galaxies experience cyclical episodes of gas accretion, star formation and black hole growth, followed by the production of a radio jet that shuts down further activity. The behaviour of galaxies with log $M_*/M_\\odot > 11.4$ is the same at $z$ = 0.6 as it is at $z$ = 0.2, except that higher redshift galaxies experience more star formation and black hole growth and produce more luminous radio jets during each accretion cycle. ",
        "introduction": "\\label{sec:intro} Large redshifts surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{york00} and the Two Degree Field (2DF) redshift survey \\citep{colless01} have provided the sky coverage needed to define sufficiently large samples of nearby radio sources to study the population statistics and global energetics of these systems at $z \\sim 0.1$. \\citet{best05b} investigated the properties of a sample of 2215 nearby radio galaxies with 1.4 GHz radio luminosity below $10^{25}{\\rm W~Hz}^{-1}$ from SDSS, finding that the fraction of radio-loud AGN is a strong function of both black hole and stellar mass. \\citet{mauch07} studied a sample of 2661 radio-loud AGN selected from the 6dF Galaxy Survey. They confirmed the findings of \\citet{best05b} that radio-loud AGN preferentially inhabit the brightest and most massive host galaxies, and showed that the fraction of galaxies which host a radio-loud AGN correlates with the infrared $K$-band luminosity as $L_K^2$. \\citet{best05b} also reported  that there was no correlation between radio luminosity and optical emission-line luminosity for the galaxies in their sample, concluding that optical AGN and low-luminosity radio-loud AGN are independent phenomena which are triggered by different physical mechanisms. In later work, it was found that at fixed stellar mass, radio-loud AGN are preferentially located in denser environments than control samples of radio-quiet galaxies with similar masses and redshifts \\citep{best07, wake08a, wake08b, mandelbaum09, donoso10} as well as control samples of optically-identified AGN \\citep{kauffmann08, donoso10}. These results have led to the paradigm that strong  optical AGN activity is associated with galaxies with a significant cold gas reservoir and ongoing star formation \\citep{heckman04}, while radio AGN activity may be associated with the accretion of hot gas in galaxies located at the centers of  massive dark matter halos \\citep{bower06, croton06}.  Observational support for a scenario in which radio AGN play an important role in regulating the growth of massive galaxies has accumulated rapidly over the past decade. X-ray studies of groups and clusters of galaxies with Chandra and {\\it XMM-Newton} have shown that these jets interact strongly with their environment, blowing clear cavities or `bubbles' in the surrounding X-ray-emitting gas \\citep{bohringer93, fabian00, churazov01, fabian03, fabian05, forman05}. These cavities provide a direct estimate of the radio jet power \\citep{birzan04, birzan08} and hence an empirical calibration of the relation between radio luminosity and mechanical power of the radio jet in nearby galaxies. The energy input from jets has been found to be sufficient to prevent star formation activity in massive early-type galaxies by heating the interstellar gas and suppressing the onset of cooling flows \\citep{binney95, birzan04, rawlings04, best06, best07, schawinski06}.  If radio jets regulate ongoing star formation in massive galaxies, one might expect to see differences in the stellar populations of the host galaxies of radio-loud AGN compared to radio-quiet galaxies. \\citet{best05b} compared the fraction of radio-loud AGN among galaxies with different values of the 4000 \\AA\\ break strength (D4000), which is a powerful measure of the mean age of the stellar population in a galaxy. As can be seen from Figure~10 of that paper, there appears to be a lower fraction of radio-loud AGN among the most massive galaxies ($> 10^{11} M_\\odot$) with young stellar populations (i.e. low values of D4000). The statistics were rather poor, so Best et~al. did not place much emphasis on this result.  While much of our knowledge about radio AGN is based upon studies of local galaxies, our understanding of the interplay between gas cooling processes and radio AGN feedback at higher redshifts is still very much in its infancy. Strong cosmological evolution of the high-luminosity radio source population has been long established  \\citep[e.g.,][]{dunlop90}, but the role of this strongly evolving population in regulating the growth of massive galaxies at high redshifts is largely unconstrained. \\citet{johnston08} constructed high signal-to-noise ratio (S/N) composite optical spectra of 375 radio-detected luminous red galaxies  at $0.4 < z < 0.8$. They compared composite spectra binned by radio luminosity and redshift  with those of radio-quiet control galaxies and concluded that there were no differences in the stellar populations of the two samples, except for the most powerful radio galaxies (with $L_{\\rm 1.4GHz}  > 10^{26}{\\rm W~Hz}^{-1}$), which had  stronger optical emission lines and  younger stellar populations. We note, however, that the galaxy survey  on which this analysis was based was  subject to a strong pre-selection on galaxy colour and would miss blue, star-forming galaxies enitirely.  The Baryon Oscillation Spectroscopic Survey (BOSS; Dawson et al. 2012, submitted), one of the surveys of SDSS-III \\citep{eisenstein11}, is currently obtaining spectra of 1.5 million luminous galaxies to $z \\sim 0.7$. At redshifts greater than $z \\sim 0.55$, the galaxy samples are reasonably complete at stellar masses greater than $10^{11.4} M_{\\odot}$. In this paper, we have cross-correlated galaxies with $z>0.55$ from the  BOSS spectroscopic survey with the National Radio Astronomy Observatory (NRAO) Very Large Array (VLA) Sky Survey \\citep[NVSS;][]{condon98} and the Faint Images of the Radio Sky at Twenty centimeters (FIRST) survey \\citep{becker95} and  have generated the largest sample of intermediate redshift radio galaxies to date. We use this sample to investigate whether radio-loud AGN have had unusual star formation histories (SFHs) compared to galaxies of similar mass and redshift. Similar to \\citet{johnston08}, we create composite optical spectra to study the emission line properties of radio-loud and radio-quiet galaxies. By comparing results obtained for the BOSS samples with those from the low-redshift radio AGN catalogue of \\citet{best12}, we study how the radio AGN population has evolved over a timescale of 3.3 Gyr.  Our paper is arranged as follows. In \\S2, we describe the low and high redshift radio-loud AGN samples and the construction of the radio-quiet control samples.  In \\S3, we discuss how we extract information about the past star formation histories of the galaxies in our samples, present results for the radio-loud AGN hosts and the control galaxies, and analyze the emission line properties of radio galaxies using composite spectra . Our  results are summarized in \\S4. We use the cosmological parameters $H_0=70~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$ throughout this paper.  \\begin{figure*} \\bc \\hspace{-0.6cm} \\resizebox{17cm}{!}{\\includegraphics{./Fig/fig_mass_lradio_dist.ps}}\\\\% \\caption{The distribution of the stellar mass and radio luminosity for the DR7 (red) and CMASS (black) radio-loud AGN samples. Our DR7 sample includes 5818 galaxies with $0.1 <  z < 0.3$, while the CMASS sample includes 2856 galaxies with $0.55 < z < 0.7$. \\label{mass_lradio_dist}} \\ec \\end{figure*} ",
        "conclusions": "We have cross-matched the SDSS-III BOSS galaxy spectroscopic sample with the NVSS and FIRST surveys, generating the largest radio galaxy sample with optical spectroscopy at $z\\sim0.6$ to date. Combining this sample with the SDSS DR7 radio-loud AGN catalogue of  \\citet{best12}, we study the evolution of the recent star formation histories of the host galaxies of radio-loud AGN, as well as  radio-quiet ``control'' galaxies that are matched in redshift, stellar mass, and velocity dispersion.  In previous work (Chen et al 2012), we found that the fraction of massive galaxies that have formed a significant fraction ($>$ a few  percent) of their stars in the past Gyr has  strongly decreased with redshift from $z \\sim 0.6$ to the present day. The goal of this paper is to analyze if this evolution is linked to the radio-emitting jets hosted by massive galaxies. Our main empirical results may be summarized as follows:  \\begin{enumerate} \\item  For galaxies with stellar masses $M_* > 10^{11.4}M_\\odot$, the fraction of actively star-forming galaxies is $\\sim 2$ times lower among galaxies with radio-emitting jets than in the  radio-quiet control samples. This is true  both at low- and at high- redshift. The factor of two difference in the fraction of star-forming galaxies is independent of  radio luminosity, except for radio galaxies with luminosities in excess of $10^{25.5}$ W Hz$^{-1}$, where the difference disappears.  \\item The locus of massive galaxies in the D4000$-$H$\\delta$A plane suggests that their star formation histories are characterized by bursts rather than low-level continuous star formation.  A smaller fraction of radio-loud galaxies have undergone significant ($>$ a few percent in mass fraction) starbursts in the last $\\sim$Gyr compared to radio-quiet objects.  \\item There is a strong correlation between \\oiii\\ EQW and stellar population age (D4000) both for radio-loud AGN and for the control sample of radio-quiet galaxies. The relation between \\oiii\\ EQW and D4000 does not change with redshift.  \\item At fixed radio power, there is strong positive evolution in \\oiii\\ EQW with redshift. The \\oiii/H$\\beta$ ratio also increases at higher redshift, i.e.; radio galaxies are more LINER-like at low redshift and more Seyfert-like at high redshift. We interpret these correlations as {\\em induced trends}, which follow from the fact that all high redshift massive galaxies have more gas,  younger stellar populations and larger black hole accretion rates.  \\end{enumerate}  Our results suggest a picture in which massive galaxies experience cyclical episodes of gas accretion, star formation and black hole growth, followed by production of a radio jet that acts to shut down the star formation in the galaxy. The behaviour of galaxies with $M_* > 10^{11.4}M_\\odot$ is the same at $z=0.6$ as it is at $z=0.2$, except that higher redshift galaxies experience more black hole growth and star formation and produce more luminous radio jets during each accretion cycle.  Let us imagine a scenario in which gas is gradually accreting onto a massive galaxy. There may be an ``accumulation phase\" where the cold gas gradually builds up in a disk  with little or no star formation. After some critical gas density is reached, the disk may become unstable and lose angular momentum \\citep{salim12}. Molecular clouds form and star formation is triggered. This corresponds to the  ``burst phase\". The radio source is triggered when gas finally reaches the central black hole. As the jet expands, it pushes away the surrounding cold gas, causing the  starburst to shut down. The radio lobes expand into the surrounding intra-cluster medium, heat the ambient gas and shut down cooling onto the central galaxy. The shutdown of cooling removes the power source of the jet and the galaxy eventually returns to its original radio-quiet accumulation phase.  This picture is consistent with a population of radio galaxies with lower cold gas content and older stellar populations than radio-quiet galaxies. The strong evolution in the fraction of massive galaxies with significant star formation may  be explained by increased cooling efficiencies in massive halos at higher redshift \\citep[see Figure~2 of ][]{white91}. If the accumulation and starburst phases have roughly equal durations, this would lead to a 2:1 ratio of starbursts in the radio-quiet and radio-loud populations, as observed. Recall, however, that clustering analyses demonstrate that radio-loud AGN are preferentially located at the centers of massive dark matter halos \\citep{best07, wake08a, wake08b, mandelbaum09, donoso10}  where gas cooling and accumulation rates are presumably higher. A fair assessment of the jet/burst duty cycle requires one to control for such environment-dependent effects. However, in CMASS samples, most galaxies are classified as central galaxies, satellites are a very small fraction of the total $-$ this is because of the very high stellar mass selection (Kovac et al. 2012, in preparation).  The missing link in our analysis of the accretion/starburst/jet cycle in massive galaxies is direct observations of the gas itself. An analysis of Chandra X-ray Observatory archival data of the hot gas in 222 nearby galaxy clusters revealed that H$\\alpha$ and radio emission from the brightest cluster galaxy is much more pronounced when the the cluster's core gas entropy is low \\citep{cavagnolo08}. Systematic studies of atomic and molecular gas in massive galaxies are still limited to small samples \\citep[e.g.][]{schawinski09, catinella10, saintonge11}. Because radio galaxies are rare, large wide-field surveys will be required before samples are large enough to study how the cold interstellar medium of massive radio-loud galaxies differs from that of their radio-quiet counterparts. In the shorter term, it will  possible to stack DR7 and BOSS spectra to search for signatures of ionized gas outflows in luminous  radio-loud galaxies and in post-burst radio-quiet galaxies \\citep[e.g.][]{chen10}. This will be the subject of future work."
    },
    "1208/1208.0344.txt": {
        "abstract": "We present the global group properties of two samples of galaxy groups containing 39 high quality X-ray selected systems and 38 optically (spectroscopically) selected systems in coincident spatial regions at  0.12$<$z$<$0.79.  The total mass range of the combined sample is $\\sim 10^{12} -  5 \\times 10^{14}$ M$_{\\odot}$.  Only nine optical systems are associable with X-ray systems.  We discuss the confusion inherent in the matching of both galaxies to extended X-ray emission and of X-ray emission to already identified optical systems.  Extensive spectroscopy has been obtained and the resultant redshift catalog and group membership are provided here.  X-ray, dynamical, and total stellar masses of the groups are also derived and presented.  We explore the effects of utilizing different centers and applying three different kinds of radial cut to our systems: a constant cut of 1\\,Mpc and two r$_{\\textrm{\\tiny200}}$ cuts, one based on the velocity dispersion of the system and the other on the X-ray emission.  We find that an X-ray based r$_{\\textrm{\\tiny200}}$ results in less scatter in scaling relations and less dynamical complexity as evidenced by results of the Anderson-Darling and Dressler-Schectman tests, indicating that this radius tends to isolate the virialized part of the system.  The constant and velocity dispersion based cuts can overestimate membership and can work to inflate velocity dispersion and dynamical and stellar mass.   We find L$_{X}$-$\\sigma$ and M$_{stellar}$-L$_{X}$ scaling relations for X-ray and optically selected systems are not dissimilar. The mean fraction of mass found in stars, excluding intra-cluster light, for our systems is $\\sim$0.014 with a logarithmic standard deviation of 0.398 dex.  We also define and investigate a sample of groups which are X-ray underluminous given the total group stellar mass.  For these systems the fraction of stellar mass contributed by the most massive galaxy is typically lower than that found for the total population of groups implying that there may be less IGM contributed from the most massive member in these systems.  80\\% of 15 underluminous groups have less than 40\\% of their stellar mass in the most massive galaxy which happens in less than 1\\% of cases with samples matched in stellar mass, taken from the combined group catalog. ",
        "introduction": "The majority of galaxies in the Universe lie in galaxy groups \\citep{Eke2004a}.  Over cosmic time, groups grow hierarchically by accreting individual galaxies and smaller groups from their surrounding filamentary structure; thus, they are evolving environments.  Even within limited redshift regimes, groups are observed to have diverse properties. Local studies \\citep[e.g.][]{Zabludoff2000} reveal that their galaxy populations vary from being dominated by early (as in typical clusters) to late-type (as in the field population) galaxies.  They range from ``poor\" groups containing a relatively small number of galaxies (commonly identified via optical selection methods) to massive systems (commonly identified via X-ray emission and weak lensing).  The typical velocity dispersion within galaxy groups is comparable to the internal velocities of the galaxies they contain, making them ideal for galaxy-galaxy mergers and interactions. Therefore, groups are both important in their own right and as the predominant environment of galaxies.  Galaxy groups are not trivial to identify.  At higher redshifts, they are most easily found via the X-ray emission of their Intra-Group Medium (IGM).  X-ray surveys are biased towards selecting groups with rich IGM, and may not be typical of the dominant group population which shapes most galaxies in the Universe.  Samples selected optically may be dominated by overdensities of galaxies not yet fully virialized.  Different physical processes are likely to be active in these two regimes and thus a comparison of groups selected via these two disparate methods can illuminate these physical phenomena.  To fully understand groups as the environment in which the majority of galaxies reside and evolve requires both a significant number of groups and significant information on the galaxy group members themselves.  Wide-field surveys such as zCOSMOS and DEEP2 have identified many galaxy groups up to redshift $\\sim$1 and $\\sim$1.3 respectively \\citep{Lilly2009,Gerke2007}.  The large sample sizes which these types of surveys yield allow for the rigorous determination of global trends in group properties.  Evolution of low-mass galaxies appears to be accelerated in groups \\citep{Iovino2010} and transformation rates such as those from late to early type galaxy morphologies and from active to passive star formation activity are more than twice that in the field \\citep{Kovak2010}.  The build-up of stellar mass on the red sequence since z$\\sim$1 involves L* galaxies moving to the red sequence preferentially in groups \\citep{Cooper2007}.  The low sampling rate and bright magnitude limits of these surveys mean, however, that the majority of groups have only a few confirmed members and thus that individual systems can be difficult to examine in detail.  A complementary approach to these large volume surveys involves studying a smaller but well defined and well sampled selection of groups.  The Group Environment Evolution Collaboration (GEEC) has taken this approach, defining samples at z$\\sim$0.5 and recently extending studies up to a redshift of 1.  Intermediate redshift work has focused on optically selected groups and examined stellar masses, colors, morphologies, and star formation histories in these systems comparing to trends observed in the field \\citep{Wilman2005a,Wilman2005b, Balogh2007,Wilman2008,McGee2008,Wilman2009,Balogh2009,McGee2011}.  Our higher redshift study involves X-ray selected systems and first results show a prominent transient population, migrating from the blue cloud to the red sequence, in these groups \\citep{Balogh2011a}.  Comparing properties such as mass, X-ray luminosity and temperature, velocity dispersion, and richness via scaling relations allows us to explore the integrated properties of groups and clusters and how they relate to one another. In clusters, minimizing the scatter in these relations is a necessity in order to obtain accurate constraints on cosmological parameters.  Through large, uniform samples, these relations are now relatively well constrained  and seem to be very tight, even up to relatively high redshifts.   Although group samples of similar size and quality are only recently available, group scaling relations exhibit a much greater scatter due to both larger measurement errors and greater intrinsic scatter in group properties \\citep[e.g.][]{OsmondPonman2004,Rykoff2008,Giodini2009,Balogh2011b}.   Understanding the scatter in the relations in the group regime is a key part of illuminating the physical processes at play.  In order to study groups spanning a significant mass and evolutionary range and to compare the results obtained from two of the most common group identification methods, we have defined two different samples within the same physical area, one via optical spectroscopy and the other via X-ray emission.  In \\citealt{Finoguenov2009} (hereafter Paper I), we presented the X-ray observations of our fields and preliminary results for our sample of X-ray selected groups.  We have since finished an extensive spectroscopic campaign, significantly improving the sampling rate and depth of galaxies in our fields, and present here our full sample of X-ray and optically selected systems.  In addition to X-ray derived luminosities and masses, well constrained membership now allows us to measure velocity dispersions and dynamical masses, stellar masses, and to search for dynamical complexity in our groups.  In this paper we present our catalog of groups and explore these global group properties.  Future work will examine the galaxy populations of these groups and search for correlations with global properties.  In \\S \\ref{sec:sampledef} we describe our samples. \\S \\ref{sec:xmeas} describes the X-ray measurements of both optical and X-ray selected groups and \\S \\ref{sec:spec} details the follow-up spectroscopy of the X-ray selected systems.  NIR measurements and galaxy stellar masses are described in \\S \\ref{sec:nirstellmass}.  Global group properties including radial cuts, membership, and velocity dispersions are detailed in \\S \\ref{sec:gzmem}. X-ray and dynamical estimates of total group mass, and the total mass in stars are presented in \\S \\ref{sec:masses}.  Dynamical complexity is explored in our systems in \\S \\ref{sec:substructure} via the Dressler-Schectman (DS) and Anderson-Darling (AD) Tests. \\S \\ref{sec:lxsig} presents the L$_{X}$-$\\sigma$ relations for our samples. The `total' group masses are compared in \\S \\ref{sec:totmass}.  We discuss the stellar  and baryon content of our systems in \\S \\ref{sec:minstars} and X-ray underluminous systems in \\S \\ref{sec:underlum}.  Throughout this paper we assume a cosmology H$_{0}=75$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}=0.3$, and $\\Omega_{\\Lambda}=0.7$ unless mentioned otherwise.  %SECTION:  GROUP SAMPLE DEFINITION ",
        "conclusions": "\\label{sec:conclusions}  We have defined two group samples at 0.12$<$z$<$0.79 in the same fields, one containing 39 high quality X-ray selected systems and the other 38 optically selected systems, in order to study groups spanning a significant mass and evolutionary range. Group membership was defined and we applied three different radial cuts: two r$_{\\textrm{\\tiny200}}$ cuts (roughly approximating a virial radius) based on the X-ray emission and velocity dispersion of the systems; and a constant 1\\,Mpc cut.  Group masses were estimated from X-ray and dynamical characteristics and stellar content -- the latter two within the differing radial cuts.  Dynamical complexity and substructure was explored using the Anderson-Darling and Dressler-Schectman tests and the shape of X-ray emission. We presented the L$_{X}$-$\\sigma$ relation for our systems which is similar to that found for nearby groups and discussed the effects of centering, radial cuts and dynamical complexity/substructure in regards to outliers in this, and other scaling relations.  Best fits to this, and to L$_{X}$-M$_{stellar}$ relations for different group samples and radial cuts were presented.  Stellar mass fractions were estimated using the X-ray and dynamical mass as proxies for the group halo mass.  Finally, evidence for a population of optical systems seemingly underluminous in X-rays given their stellar and dynamical mass was discussed. Our main conclusions are as follows:\\\\  \\noindent $\\bullet$ Confusion:\\\\[3pt] Confusion exists both in matching galaxies to extended X-ray emission and matching X-ray emission to already identified optical systems.  Until X-ray spectroscopy is available to measure the redshift of the X-ray emitting gas, completely confident matching will not be possible.  Splitting systems into X-ray detected and undetected systems designates the problem, not the solution.  These difficulties in matching make cosmological studies using groups difficult.  \\\\  \\noindent $\\bullet$ Dynamical complexity:\\\\[3pt] Dynamical complexity/substructure in a system \\textit{can} work to inflate velocity dispersion and stellar mass and may explain the position of certain outliers in the scaling relations explored here.  It is important to recall that the tests we are using are orbit dependent and can only be confidently applied to systems having at least ten members.\\\\  \\noindent $\\bullet$ Radial cuts:\\\\[3pt] Applying X-ray based r$_{\\textrm{\\tiny200}}$ radial cuts usually produces the tightest scaling relations.  The good correlation between L$_{X}$ and $\\sigma$ and the lack of dynamical complexity found for systems using this radius implies that it is isolating the virialized part of the group.  Velocity dispersion based and constant cuts generally result in larger radii, more members, and include more substructure/non-Gaussianity.  This acts to increase scatter and inflate both velocity dispersion and stellar mass. However, as some systems are not X-ray detected, such cuts are the only options.\\\\  \\noindent $\\bullet$ Stellar mass fraction:\\\\[3pt] We find a mean stellar mass fraction of $\\sim$0.014 within an X-ray based r$_{\\textrm{\\tiny200}}$ and treating the X-ray mass as the total mass of the system.  This is comparable to those found by \\cite{Giodini2009}, \\cite{Balogh2011b} and \\cite{Leauthaud2011} but significantly lower than that found by \\cite{Gonzalez2007}.  The mean contribution of the most massive galaxy is $\\sim$0.004.  Using a total mass based on dynamical mass would result in different fractions due to significant disagreement between M$_{X}$ and M$_{dyn}$ for many of our systems. \\\\  \\noindent $\\bullet$ Total mass measures:\\\\[3pt] The differences in total mass measures (M$_{X}$ and M$_{dyn}$) tend to increase, and the scatter decrease, as X-ray mass increases. \\\\  \\noindent $\\bullet$ X-ray underluminous groups:\\\\[3pt] We define a sample of systems as X-ray underluminous given their stellar mass, the majority of which are optically selected.  Not all such systems show dynamical complexity and the stellar mass fraction in the most massive galaxy of these systems is on average less than that found for the total population of groups.   This may indicate that less IGM is being contributed from the progenitor halo containing the most massive member and we posit that differences in accretion (a continuous smooth accretion of galaxies from the field verses the merging of similar mass `subgroups')  may be one explanation for the observed correlation between the fraction of mass in the most massive galaxy and the relative X-ray luminosity. \\\\"
    },
    "1208/1208.1856_arXiv.txt": {
        "abstract": "We examine a sample of 48 {\\it Swift}/UVOT long Gamma-ray Burst light curves and find a correlation between the logarithmic luminosity at 200s and average decay rate determined from 200s onwards, with a Spearman rank coefficient of -0.58 at a significance of 99.998\\% (4.2$\\sigma$). We discuss the causes of the $\\rm log\\;L_{200\\rm{s}}-\\alpha_{>200\\rm{s}}$ correlation, finding it to be an intrinsic property of long GRBs, and not resulting from the selection criteria. We find two ways to produce the correlation. One possibility is that there is some property of the central engine, outflow or external medium that affects the rate of energy release so that the bright afterglows release their energy more quickly and decay faster than the fainter afterglows. Alternatively, the correlation may be produced by variation of the observers viewing angle, with observers at large viewing angles observing fainter and slower decaying light curves. ",
        "introduction": "\\label{intro}  Gamma-ray bursts (GRBs) are intense flashes of gamma-rays that are usually accompanied by an afterglow, longer lived emission that may be detected at X-ray to radio wavelengths. Our understanding of GRB X-ray and optical/UV afterglows was revolutionised by the launch of {\\it Swift}, a satellite dedicated to the detection of GRBs and observation of their $\\gamma$-ray, X-ray and optical/UV emission \\citep{geh04}. The best studied subclass of GRB are the long GRBs (LGRBs), which have observed prompt $\\gamma$-ray emission durations of $\\gtrsim 2$s \\citep{kou93}. Several studies have investigated the X-ray emission of LGRBs, using large samples, to identify characteristic temporal and spectral behaviours \\citep[i.e][]{nousek,eva07,eva09}. Similar investigations have been performed at optical/IR wavelengths \\citep[i.e][]{mel08,ryk09,oates09}, but these have tended to have much smaller samples, due to the lower detection rate of the optical emission in comparison to the X-rays (the detection rate for {\\it Swift}'s X-ray and ultra-violet optical telescopes are $\\sim96\\%$ \\citep{bur08} and $\\sim29\\%$ \\citep{rom09}, respectively for LGRBs). The low detection rate is generally attributed to extinction due to high levels of dust in the host \\citep{fyn01}, and/or to high redshift at which the optical/UV emission will be absorbed by neutral hydrogen along the line of sight \\citep{gro98}; see also \\cite{gre11} for a recent study on the cause of optically dark GRBs. However, {\\it Swift} has now observed over 600 GRBs and we are now in a position to collate a large number of well-sampled IR/optical/UV LGRB afterglows from which we can draw a representative picture of their behaviour and collective properties \\citep[see recent papers by][]{kan10,li12}.  In \\cite{oates09}, we performed a statistical investigation of 26 optical/UV LGRB afterglows, finding a correlation between the observed $v$-band magnitude at 400s and the average UVOT light curve decay rate determined from 500s. We also tested for an equivalent rest frame correlation, but, due to the small sample size, this could not be confirmed or excluded. Here we use a larger sample of 48 high quality LGRB UVOT light curves to re-examine if there is a correlation between optical/UV afterglow intrinsic brightness and light curve decay rate.  This paper is organized as follows. In \\S~\\ref{reduction} we describe the data reduction and analysis. The main results are presented in \\S~\\ref{results}. Discussion and conclusions follow in \\S~\\ref{discussion} and \\ref{conclusions}, respectively. All uncertainties throughout this paper are quoted at 1$\\sigma$. The temporal and spectral indices, $\\alpha$ and $\\beta$, are given by the expression $F(t,\\nu)\\propto t^{\\alpha}\\nu^{\\beta}$.  \\begin{figure} \\includegraphics[angle=-90,scale=0.30]{Lum_light_curves.cps} \\caption{Optical luminosity light curves of 56 GRBs at restframe 1600\\AA. For clarity, 3 $\\sigma$ upper limits are not included.} \\label{lightcurves} \\end{figure} ",
        "conclusions": "\\label{conclusions} We computed luminosity light curves at 1600{\\AA} for 48 optical/UV GRB afterglows. We find a correlation between luminosity at 200s and average decay rate from 200s onwards with a significance of 99.998\\% (4.2$\\sigma$). Regression analysis indicates a linear relationship between decay rate and luminosity of $\\rm log\\;L_{200\\rm{s}}=(-3.636\\pm0.004)\\alpha+(28.08\\pm0.13)$.  We used a Monte Carlo simulation to determine, at 4.1$\\sigma$ confidence, that the $\\rm L_{200\\rm{s}}-\\alpha_{>200\\rm{s}}$ correlation is intrinsic and not due to chance or our selection criteria. We determined that this correlation is not likely to be a natural consequence of the basic synchrotron afterglow model. Instead we find two possible ways to produce the correlation. The first is that there is some property of the central engine, outflow or external medium that affects the rate of energy release and rate of light curve decay, in such a way that for brighter afterglows the energy is released more quickly and decays more rapidly than the fainter afterglows. Alternatively, the correlation may be produced by a range in observing angles, with observers at large viewing angles witnessing fainter and slower decaying light curves. Understanding the origin of this correlation will have important consequences on our understanding of the physics and geometry behind GRBs.  \\vspace{-1mm}"
    },
    "1208/1208.5472_arXiv.txt": {
        "abstract": "We present a model of polarization swings in blazars from axially symmetric blobs propagating on curved trajectories. If the minimum inclination of the velocity vector to the line of sight is smaller than $\\Gamma^{-1}$, the polarization angle maximum rotation rate is simultaneous with the polarization degree minimum and a spike in the total flux. By measuring the maximum rotation rate and the moment of the polarization maximum, it is possible to estimate the distance covered by the blob and thus its approximate position. We apply this model to the recent polarization event in blazar 3C~279. ",
        "introduction": "Radio to optical emission of relativistic jets in Active Galactic Nuclei (AGNs) is linearly polarized and this fact gives the strongest support for its synchrotron origin. Both polarization degree (PD) and polarization angle (PA) are strongly variable. Most of the time their behaviour is rather chaotic, but occasionally they show coherent events. Two kinds of PA variations have been observed: systematic rotations of amplitude much larger than $180^\\circ$ on timescales of months to years ({\\it e.~g.} in PKS 0721-115\\cite{1981ApJ...248L...5A} and PKS 0521-365\\cite{1993A&A...269...77L}) and fast swings of amplitude close to $180^\\circ$ ({\\it e.~g.} in S5 0917+624\\cite{1989A&A...226L...1Q} and BL Lac\\cite{2008Natur.452..966M}).  Theoretical efforts to explain the PA rotation in AGN jets began with a model of accelerating clumps of matter\\cite{1979ApJ...232...34B}.  A net transverse magnetic field component is required to break the axial symmetry of the emitting region, so that the resulting PA rotation is smooth and gradual. It was later shown that other kinematic effects causing variations of the viewing angle in comoving frame (jet bending or reorientation) could trigger a PA swing\\cite{1982ApJ...260..855B}. These models could explain PA swings of amplitude up to $180^\\circ$. For larger PA rotations, models involving helical\\cite{1985ApJ...289..188K} or stochastic\\cite{1985ApJ...290..627J} magnetic fields have been proposed.  In all these early works some kind of non-axisymmetric internal jet structure was assumed. Here we show, that it does not need to be the case. A symmetric emitting region propagating on a bent trajectory, not necessarily helical, can produce a gradual PA swing. This is arguably the simplest way to reproduce several features in the simultaneous behaviour of PA and PD in some observational datasets with a minimum number of parameters. ",
        "conclusions": "Our model is most suitable for PA swings of $\\sim 180^\\circ$ with simultaneous PD minimum and brightness spike, and with characteristic fast rise and slow decay of PD. We note that reported amplitude $208^\\circ$ of the polarization rotation in 3C~279 is based on the values measured immediately before the event, when there was a small rotation in the other direction. We also predict an apparent $\\sim 180^\\circ$ twist in a pc-scale jet, that in general could be verified with high-resolution VLBI observations. However, in the case of 3C~279 this is not plausible, since there is little variability in the mm wavelenghts and the optical/$\\gamma$-ray flare is produced within the mm photosphere.  Larger observed amplitudes of PA rotation, {\\it e.~g.} $\\sim 240^\\circ$ in BL Lac\\cite{2008Natur.452..966M}, could be explained by helical trajectories. If the emission source is located in the Poynting-flux-dominated region, they would be determined by conservation of angular momentum inherited from the central engine. However, they could also arise in the interactions of the jet with non-uniform and variable environment on several-pc scales. There are two observational facts supporting this scenario in the case of 3C~279. First, the kpc-scale jet of continuously changes its positional angle and the apparent velocity of the radio knots on timescale of years\\cite{2008ApJ...689...79C}. Second, a PA rotation in the other direction has been observed in earlier epoch\\cite{2008A&A...492..389L}, this challenges the models predicting fixed helicity sign.  If the emitting region is large enough with respect to the curvature radius, time delays between the arrivals of photons emitted from different positions across the emitting volume would affect the variation of polarization parameters. This would provide a practical method to estimate the radius of the emitting region. We could also be able to place a constraint on the magnetic field distribution across the blob extension, in particular distinguish between globally ordered (toroidal) and disordered (chaotic) magnetic fields.  Long and intensive polarimetric monitoring of bright blazars is very important, since polarization swings are fast and rare events, that cannot be predicted in advance. A good coverage is essential for verifying competing theoretical models. In the case of 3C~279 an important question is, whether the observed polarization swing was related to the transition from high to low $\\gamma$-ray flux, or was it purely coincidental. We will only know this, when we observe another polarization swing in similar circumstances."
    },
    "1208/1208.5770.txt": {
        "abstract": "The observed abundance of giant arcs produced by galaxy cluster lenses and the measured Einstein radii have presented a source of tension for \\lcdm, particularly at low redshifts ($z\\sim0.2$). Previous cosmological tests for high-redshift clusters ($z>0.5$) have suffered from small number statistics in the simulated sample and the implementation of baryonic physics is likely to affect the outcome. We analyse zoomed-in simulations of a fairly large sample of cluster-sized objects, with $\\mvir>3\\times10^{14}\\,h^{-1}\\msol$, identified at $z=0.25$ and $z=0.5$, for a concordance $\\Lambda$CDM cosmology. These simulations have been carried out by incrementally increasing the physics considered. We start with dark matter only simulations, and then add gas hydrodynamics, with different treatments of baryonic processes: non\u00d0radiative cooling, radiative cooling with star formation and galactic winds powered by supernova explosions, and finally including the effect of AGN feedback. Our analysis of strong lensing properties is based on the compuation of the cross-section for the formation of giant arcs and of the Einstein radii. We find that the addition of gas in non--radiative simulations does not change the strong lensing predictions significantly, but gas cooling and star formation together significantly increase the number of expected giant arcs and the Einstein radii, particularly for lower redshift clusters {\\it and} lower source redshifts. Further inclusion of AGN feedback reduces the predicted strong lensing efficiencies such that the lensing probability distributions becomes closer to those obtained for simulations including only dark matter. Our results indicate that the inclusion of baryonic physics in simulations will not solve the arc-statistics problem at low redshifts, when the physical processes included provide a realistic description of cooling in the central regions of galaxy clusters. As outcomes of our analysis, we encourage the adoption of Einstein radii as a robust measure of strong lensing efficiency, and provide the \\lcdm~predictions to be used for future comparisons with high-redshift cluster samples. ",
        "introduction": "The hierarchical formation paradigm describes a scenario in which small structures form earlier and then merge to form more massive objects. This description within the \\lcdm~model has been highly successful at explaining observations of structure on a large range of scales and redshifts \\citep[e.g.][]{ESM90,WF91,K11}, including the formation of galaxy clusters \\citep[e.g.][]{KB12}. However, the internal structure of galaxy clusters has posed a challenge. The earliest comparisons between simulated clusters and the observed frequency of gravitational lensing arcs revealed a serious discrepancy between the observations and \\lcdm~predictions \\citep{B98,L05}. More recent comparisons to X-ray selected clusters find that the discrepancy remains for $z<0.3$ and $z>0.5$ but results are not definitive due to small simulated sample sizes at higher redshifts and uncertainties with selection procedures and biases \\citep{M11,H11}. In addition, observations of clusters at redshifts $0.1<z<0.6$ have found higher concentrations and larger Einstein radii than predicted \\citep{BB08, SR08, Z11a}. These are related problems since both suggest a failure of \\lcdm~to correctly infer the distribution of matter in the cores of clusters; the solution may be found in the details of the simulations used to formulate the predictions.  Simulations which only incorporate collisionless particles are commonplace in the literature. Dark matter structure is evolved under gravity with the effect of baryons ignored. On large-scales, the distribution of gas follows the dark matter potential wells, so the collisionless simulations will suffice; however on galaxy- and group- scales, gas cooling, star formation and feedback can alter the gravitational potential significantly \\citep[e.g.][]{BFFP86,G04}. Early simulations that were used to model the cluster lenses did not incorporate sufficient relevant physical processes \\citep{HWY07,SR08}. These simulations were originally justified by the premise that the temperature of the ICM is too hot to allow efficient cooling, so the baryonic matter must simply follow the dark matter potential and would not have a significant impact on its shape \\citep[e.g.][]{BB08}. The role of baryons in shaping galaxy clusters has recently become a key point of discussion since hydrodynamic simulations have become more feasible; the findings are that baryons do, in fact, affect the shapes of gravitational potential of the simulated haloes. Gas cooling, for example, leads to adiabatic compression of galaxy haloes preferentially in the cluster centre, allowing more mass to be deposited around the core \\citep{BL10}. Cooling and star formation serve to steepen the density profile of clusters and thus increase the mass concentration and the strong lensing cross-section \\citep{Lewis00,G04,P05,R08,RZK08,Cui12}. Unfortunately, these processes lead to the well known overcooling problem, in which the cold gas mass and stellar mass in the cluster core is overestimated.  %Why AGN Feedback is important Including in simulations feedback from gas accretion onto supermassive black holes (SMBHs) is a promising solution for overcooling \\citep{Sijacki07,T11,McC10,Fabjan10}, while simultaneously reproducing the drop in the cosmic SFR for $z<2$ \\citep[e.g.,][]{vdV11}. The resulting feedback from active galactic nuclei (AGN) significantly reduces the gas fraction in simulated galaxy-groups and poor clusters ($T\\lesssim2$keV), but cannot remove gas from the deep potential wells of rich, massive clusters \\citep{PSS08,Fabjan10,McC10}. While the total baryonic content of massive clusters is relatively unaffected by cooling, star formation and energy feedback, such processes can imply a quite pronounced redistribution of baryons, particularly close to the cluster centre. \\citet{D10} analysed simulated haloes at $z=0$ ranging from galaxy to cluster scales. They demonstrated that while radiative cooling increases mass concentration\\citep[see also][]{RZK08}, AGN feedback has the opposite effect; clusters modelled with AGN feedback have mass concentrations equal to those modelled with dark matter only. Consistent with these results, \\citet{Mead10} demonstrated that AGN feedback is able to reduce the strong lensing cross section for clusters at $z=0.2$. Four of the five clusters they simulated had strong lensing efficiencies consistent with their dark-matter counterparts, while the fifth remained a stronger lens.  % Drawbacks of previous studies Many of the previous works investigating the role of baryons in modifying cluster cores have told similar stories: while cooling and star-formation either directly or indirectly leads to an increase of mass in the core, and therefore an increase in the strong lensing efficiency, AGN feedback negates this to some extent. Unfortunately previous studies have had limited samples of massive clusters, especially at high-redshift $z\\gtrsim0.3$, and have utilised simulations generated with values for cosmological parameters that have since been revised. The strong lensing studies of \\citet{P05}, \\citet{R08} and \\citet{Mead10} were undertaken in the WMAP-1 type cosmology with $\\sig=0.9$. As shown by \\citet{M08}, changing the cosmology from WMAP-1 to WMAP-3 best fitting models can lead to a 20\\% decrease in the predicted mass concentration of relaxed clusters. This is primarily due to the decrease of the power-spectrum normalisation, $\\sig$, from 0.9 to 0.7, and the subsequent delay in the assembly of haloes. The current WMAP-7 best-fit values lie between these two extremes; the expected reduction in the concentration of masses considered in this work from WMAP1 to WMAP5/7 is closer to 14\\% \\citep[e.g.][but see also \\citealt{Prada12}]{ M08, D08}.  % The present work In this work we improve upon the aforementioned works by analysing the strong lensing efficiencies of a larger sample of clusters simulated within the currently favoured cosmology ($\\sig=0.8$) with only dark matter (DM), and also including the hydrodynamical treatment of gas, along with cooling, star formation and feedback from both supernova (SN) and AGN. Strong lensing properties of our fairly large sample of massive clusters are analysed at two redshifts, $z=0.25$ and 0.5. A description of the cosmological simulations and the cluster sample follows in Section \\ref{sims}. Earlier studies have characterised strong lensing by the cross section for the formation of giant arcs. A comparison with observational data requires the cross-section, as a function of source redshift, to be convolved with the source redshift distribution; uncertainties in this source redshift distribution weaken the cosmological test. We, therefore, propose that the Einstein radius is a more robust proxy for strong lensing; critical curves are determined at a single source redshift (usually $z_{s}=2$), regardless of the range of source redshifts used for lens mass reconstructions. Therefore, in Section \\ref{Xsec} and Section \\ref{ER} we present the strong-lensing properties for the relaxed sub-sample and discuss the influence of the baryonic processes; in the former section we characterise strong-lensing with the cross-section for the formation of giant-arcs, while in the latter we consider the Einstein Radii instead. At the end of this section, we describe the strong-lensing properties of unrelaxed sub-sample. These results are then interpreted in terms of variation of the density profiles induced by the presence of baryons in the simulated clusters. Finally, we summarise of our findings in Section~\\ref{conclude}.   %--------------------------------------------------------------------------------------------------------------------------------- ",
        "conclusions": "%Summarize findings In the present study, we have analysed the strong lensing efficiency of about 15 relaxed simulated clusters at $z_{L}=0.25$ and $z_{L}=0.5$. The clusters were re-simulated under a number of different physical models: dark matter only ({\\tt DM}); non-radiative gas ({\\tt NR}); cooling, star-formation, SN feedback ({\\tt CSF}); and additional AGN feedback ({\\tt AGN}). Our main findings are as follows. % \\begin{itemize} \\item There are no significant differences in the strong lensing properties and density profiles of clusters modelled with dark matter only and those that include a non-radiative gas component. \\item Gas cooling and star formation have the effect of increasing the giant arc counts and Einstein radii even in the presence of stellar feedback, which can push gas out of galaxies and reduce star formation. \\item When AGN feedback is included in the simulations, the strong lensing efficiency of the clusters are significantly reduced. Einstein radii at $z_{s}=2$ are only 5 per cent larger than those for the {\\tt DM} counterparts, a boost that is increased to 10--20 per cent for lower source redshifts. There is a smaller orientation dependence, reflecting a more spherical shape due to gas cooling, even in the presence of AGN feedback. Small Einstein radii ($\\theta \\lesssim 3$\") are less likely to be produced. \\item The median Einstein radius is strongly correlated with the giant arc cross section for the relaxed cluster sub-sample. Unrelaxed clusters are especially problematic for the calculation of strong lensing efficiency using the production of giant lensing arcs as a proxy; they also have smaller strong lensing efficiencies. \\item There exists a scaling relation between strong lensing efficiencies and mass, particularly for higher overdensities ($\\Delta=2500$). As we move to lower overdensities, the normalisation for the each simulation differs primarily due to the redistribution of baryons near the cluster core. \\item We search for an explanation for these findings by analysing the mass profiles of our cluster sample; the distribution of mass in the centres of clusters are significantly influenced by baryonic physics. AGN feedback reduces the stellar mass in the core of otherwise `overcooled' clusters. Nevertheless, within $r/\\rvir \\lesssim 0.03$ the total mass is still significantly larger than that of the {\\tt DM} counterparts. However, this does not translate to higher lensing efficiencies for $z_{s}=2$, because $\\EinRad \\gtrsim 0.03 \\rvir$ for most clusters in our chosen mass range. \\end{itemize}  %Discussion about choice of implementation and comparison to previous works Where there is overlap in cluster masses and redshifts, our results are similar to those obtained by \\citet{D10} and \\citet{Mead10} with regards to the effect of AGN feedback on density profiles and giant-arc cross-sections. The AGN feedback prescription used in each work are all inspired by the model of \\citet{SdMH05} but with somewhat different implementations and choice of the relevant model parameters. Therefore, it is reassuring that the strong lensing predictions are not sensitive to these details. However \\citet{T11}, who include mechanical AGN feedback associated with jets into an Eulerian AMR code, have noted that stellar fractions are sensitive to implementation. An open question is then whether different implementations of AGN feedback in different codes produce comparable effects on the inner density profiles of massive halos and, therefore, on the strong lensing properties of galaxy clusters.  % Comparison to Observations As for the comparison with observational data, available results on strong lensing arc statistics for samples of nearby and distant clusters have been recently carried out, for instance, by \\cite{H10}. Furthermore, the ongoing Cluster Lensing And Supernova survey with Hubble \\citep[CLASH;][]{P12} project, thanks to the superior sensitivity of deep HST imaging, is providing unprecedented detail in the internal mass distribution of galaxy clusters through strong lensing studies. These results represent an excellent test ground for the comparison with simulation results, like those presented in this paper. Self-consistent comparisons with data are deferred for future study, as this will require more careful selection of clusters, with a comparable mix of relaxed and unrelaxed objects in the observational and simulated samples, as well as consistent methods of measuring strong lensing properties. As \\citet{WTS01} have shown, a reduction in the lensing efficiency due to a decreased mass concentration may be partially compensated by a higher magnification for individual images. Therefore, a comparison of giant arc counts will require not only a prediction for the strong lensing cross-section, but also the magnification of images that are {\\it potential} giant arcs, in order to determine if they would be detected in flux-limited imaging; this is particularly important for images that lie near the flux limit.   % Final statement Approaches to resolving the arc statistics problem have generally involved questioning three aspects of the cosmological test: the assumed source redshift distribution, the realism of simulations, and the self-consistency in selection criteria for simulated and observed clusters. We prefer the Einstein radius as a proxy for strong lensing, since it allows one to avoid uncertainties with describing realistic source populations. None of the simulated models are tuned to reproduce {\\it all} observable properties of clusters, but the {\\tt AGN} set of simulations is the most realistic one. Our study shows that AGN feedback is a game-changer; it's hard to explain a possible discrepancy between observed and simulated lensing efficiency by resorting to the effect of baryons. Therefore, it will be crucial to make self-consistent comparisons between simulations and observations at a range of redshifts to understand if lensing efficiency truly presents a challenge for \\lcdm.    %---------------------------------------------------------------------------------------------------------------------------------"
    },
    "1208/1208.4124_arXiv.txt": {
        "abstract": "We present infrared observations of the ultra-compact H {\\small II} region W3(OH) made by the FORCAST instrument aboard SOFIA and by Spitzer/IRAC. We contribute new wavelength data to the spectral energy distribution, which constrains the optical depth, grain size distribution, and temperature gradient of the dusty shell surrounding the H {\\small II} region. We model the dust component as a spherical shell containing an inner cavity with radius $\\sim 600$ AU, irradiated by a central star of type O9 and temperature $\\sim 31,000$ K. The total luminosity of this system is $7.1 \\times 10^4~L_\\odot$. An observed excess of $2.2 - 4.5~\\mu$m emission in the SED can be explained by our viewing a cavity opening or clumpiness in the shell structure whereby radiation from the warm interior of the shell can escape. We claim to detect the nearby water maser source W3 (H$_2$O) at 31.4 and $37.1~\\mu$m using beam deconvolution of the FORCAST images. We constrain the flux densities of this object at $19.7 - 37.1~\\mu$m. Additionally, we present {\\it in situ} observations of four young stellar and protostellar objects in the SOFIA field, presumably associated with the W3 molecular cloud. Results from the model SED fitting tool of \\citet{Robitaille06,Robitaille07} suggest that two objects (2MASS J02270352+6152357 and 2MASS J02270824+6152281) are intermediate-luminosity ($\\sim 236 - 432~L_\\odot$)  protostars; one object (2MASS J02270887+6152344) is either a high-mass protostar with luminosity $3\\times10^3$ L$_{\\odot}$ or a less massive young star with a substantial circumstellar disk but depleted envelope; and one object (2MASS J02270743+6152281) is an intermediate-luminosity ($\\sim 768~L_\\odot$) protostar nearing the end of its envelope accretion phase or a young star surrounded by a circumstellar disk with no appreciable circumstellar envelope. ",
        "introduction": "Ultra-compact H {\\small II} regions can be found throughout the galaxy surrounding newly formed massive stars just completing the accretion stage and entering their main sequence lifetimes. UCHs are small ($<0.4$~pc), hot ($>100$~K), and massive ($>10^{3} M_{\\odot}$), and emit $10^{4} - 10^{6}~L_\\odot$~ \\citep{Churchwell}. These regions of ionized gas are surrounded by molecular and dust clouds, out to as far as ten times the radius of the UCHs themselves \\citep{Conti}. This causes attenuation of much of the emitted luminosity, as it is absorbed and reradiated in the infrared.  W3(OH) is one of the largest and best studied UCH {\\small II} regions known in the Galaxy. It is at a distance of approximately 2.04 kpc in the Perseus arm \\citep{Hachisuka}. Like most UCH \\small{II} regions, is surrounded by its natal dust and molecular gas envelope as well as a larger giant molecular cloud encompassing W3(OH) and W3(H$_{2}$O), W3 Main, and AFGL 333; these are all regions of potentially triggered star formation, based on their positions on the dense outskirts of a much less dense cavity in the W3 GMC \\citep{Ruch, Moore}. Further review of star formation in the W3 GMC is presented by \\citet{Megeath}.  Line emission from the molecular gas surrounding W3(OH) has been mapped at submillimeter and radio wavelengths in the transitions of OH and H$_2$O \\citep{mader78},  HCN \\citep{Turner}, NH$_3$  \\citep{wilson78, zeng84, tieftrunk98}, CH$_3$OH \\citep{Menten}, C$^{18}$O \\citep{Wink}, and C$^{17}$O \\citep{Wyrowski97}. This molecular emission is found out to a $1'$ diameter region. These studies have particularly focused on further characterizing the molecular gas and the OH and methanol masers within and around W3(OH), as well as the hot core/water maser region W3(H$_{2}$O) approximately 6'' to the east.  In continuum emission, however, the dust cocoon surrounding the UCH {\\small II} region comes into primary focus.  The dust component to the UCH {\\small II} W3(OH) was detected by \\citet{wynnwilliams72} at $\\sim 1 - 20~\\mu$m; the dust component is optically thick in the near- and mid-infrared. \\citet{chini86} studied the cold dust at wavelengths of 350 $\\mu$m and 1.3 mm. Using a spherically symmetric radiative transfer model, they concluded that the inner cavity ($\\lesssim 2 \\times 10^{17}$~cm) of the UCH {\\small II} region is depleted of dust, rather than having dust density that increases approaching the central star. A large amount of visual extinction (A$_{\\rm v} \\sim 67$) in a thick outer shell was required to explain the decline in emission at $\\lambda \\lesssim 5~\\mu$m. Utilizing airborne data, \\citet{campbell89} showed that the dust cocoon is optically thick in the far-infrared and developed a more detailed model containing an H {\\small II} region and a cavity. Surrounding the cavity is a dusty region with a free-fall density distribution and a temperature gradient through the dust cloud. More recently, \\citet{Stecklum} presented high spatial resolution, ground-based 10 and 20 $\\mu$m images of W3(OH). From these data, the model for the dust shell was further developed and contained a Gaussian density distribution with an inner cavity of radius 2270 AU and stellar luminosity of $8 \\times 10^4~L_\\odot$.  In this work, we present new, high spatial resolution observations of the W3(OH) region in the wavelength range $\\sim 3.6 - 40~\\mu$m obtained with the FORCAST instrument \\citep{herter12} on SOFIA \\citep{young12} and with the Infrared Array Camera (IRAC; Fazio et al. \\citeyear{fazio04}) on the Spitzer Space Telescope \\citep{werner04}. These wavelengths are critical for determining the spectral energy distribution (SED) of the W3(OH) dust cocoon and thus for making a measurement of its total luminosity. We combine our data with 2MASS \\citep{Skrutskie06} fluxes and other data published in the literature to construct SEDs for the W3(OH) dust component. We model the dust component as a dusty shell around the H {\\small II} region and compute the emergent flux using radiative transfer code. In addition, we present {\\it in situ} observations of four stellar and protostellar objects in the SOFIA/FORCAST field. We fit the SEDs of these objects to those of high- and intermediate-mass protostars containing dusty circumstellar envelopes and circumstellar disks. We discuss the physical properties of all these objects. ",
        "conclusions": "We present SOFIA/FORCAST and Spitzer/IRAC observations of the UCH {\\small II} region W3(OH) in the wavelength range $3.6 - 37.1~\\mu$m. These data, combined with other published data, have been used to constrain the optical depth, grain size distribution, and temperature gradient in the dusty shell surrounding the H {\\small II} region. The total luminosity of W3(OH) is $7.1 \\times 10^4~L_\\odot$, indicating that the central star is an O9 star with surface temperature $\\sim 31,000$ K. A clumpy dust distribution or cavity opening revealing warm interior grains is necessary to explain excess emission at $2.2 - 4.5~\\mu$m.  We detect the hot core W3(H$_2$O) at 31.4 and $37.1~\\mu$m, and constrain its flux density at $19.7 - 37.1~\\mu$m using deconvolved FORCAST images.  In addition, SEDs have been constructed for four young stellar or protostellar objects which lie in the SOFIA/FORCAST field. The model SED fitting tool of \\citet{Robitaille06} was used to determine the nature of these objects. 2MASS J02270352+6152357 is an intermediate-luminosity protostar undergoing envelope accretion; 2MASS J02270824+6152281 is most likely a very young intermediate-mass protostar with a large natal envelope; 2MASS J02270887+6152344 is a high-luminosity object which is either a protostar with ongoing envelope accretion onto a young disk or a young star with a circumstellar disk and a depleted envelope; and 2MASS J02270743+6152281 could be an intermediate-luminosity protostar or potentially a young star with a developed disk and an almost entirely depleted envelope.  Further observations in the mid-IR, far-IR and/or submillimeter range(s) are required to definitively characterize 2MASS J02270887+6152344  and 2MASS J02270743+6152281."
    },
    "1208/1208.3368_arXiv.txt": {
        "abstract": "We report the first investigation of cool-core properties of galaxy clusters selected via their Sunyaev--Zel'dovich (SZ) effect.  We use 13 galaxy clusters uniformly selected from 178 deg$^2$ observed with the South Pole Telescope (SPT) and followed up by the {\\it Chandra X-ray Observatory}.  They form an approximately mass-limited sample ($> 3\\times 10^{14} \\Msun h^{-1}_{70}$) spanning redshifts $0.3 < z < 1.1$.  Using previously published X-ray-selected cluster samples, we compare two proxies of cool-core strength: surface brightness concentration (\\csb) and cuspiness ($\\alpha$).  We find that \\csb\\ is better constrained.  We measure \\csb\\ for the SPT sample and find several new $z > 0.5$ cool-core clusters, including two strong cool cores.  This rules out the hypothesis that there are no $z > 0.5$ clusters that qualify as strong cool cores at the 5.4$\\sigma$ level.  The fraction of strong cool-core clusters in the SPT sample in this redshift regime is between 7\\% and 56\\% (95\\% confidence).  Although the SPT selection function is significantly different from the X-ray samples, the high-$z$ \\csb\\ distribution for the SPT sample is statistically consistent with that of X-ray-selected samples at both low and high redshifts.  The cool-core strength is inversely correlated with the offset between the brightest cluster galaxy and the X-ray centroid, providing evidence that the dynamical state affects the cool-core strength of the cluster.  Larger SZ-selected samples will be crucial in understanding the evolution of cluster cool cores over cosmic time. ",
        "introduction": "Galaxy clusters grow over cosmic time through mergers with other galaxy clusters as well as through the accretion of gas and individual galaxies from the surrounding environment.  On timescales of a few Gyr, radiative cooling due to X-ray emission from the intracluster medium (ICM) would give rise to a ``cooling flow'' to the cluster core \\citep{fabian77,mathews78}, if it were not countered by a heating mechanism.  These cooling flows are not observed; instead, the cores of some clusters are found to undergo only moderate cooling \\citep{kaastra01,peterson01,tamura01}.  Such galaxy clusters are known as ``cool-core'' clusters \\citep{molendi01}.  Other clusters exhibit little to no cooling in their core (i.e., noncool-core clusters).  These cooling properties suggest that there must be processes in every cluster that are strong enough to either regulate cooling flows or completely prevent them.  Such processes are not fully understood and it is still uncertain how they evolve and affect cluster formation over time.  The important astrophysical processes that counteract cool-core formation are typically thought to fall under three broad categories: feedback from active galactic nuclei (AGNs), preheating of the cluster gas, and cluster mergers.  AGN feedback in the cluster's brightest cluster galaxy (BCG) has been shown to be capable of regulating cooling flows in cool-core clusters \\citep[see][for reviews]{fabian12,mcnamara12}.  In some cases, the feedback may be strong enough to disrupt a cool core completely \\citep[e.g.,][]{mcdonald11b}, though this phenomenon is likely limited to lower mass galaxy clusters and groups. Additionally, AGN feedback may drive turbulence in the ICM.  This has been shown to suppress heat-flux-driven buoyancy instabilities, resulting in effective transfer of heat from the outer radii and disrupting the cool core \\citep{parrish12b}.  Heating of the intracluster gas during early stages of the cluster has been shown to affect the formation of cool cores as well \\citep[e.g.,][]{kaiser91,mccarthy08,sun09b}.  Cluster mergers can also disrupt cool cores by shock-heating and turbulent mixing \\citep{leccardi10,rossetti11}, a process that has been reproduced in simulations \\citep[e.g][]{mcglynn84,gomez02b,zuhone10}.  Whether a merger can destroy a cool core likely depends on the strength of the cool core, the mass ratio of the merging clusters, and the geometry of the impact.  Studying the evolution of clusters can provide insight regarding the relative importance of these processes in cool-core formation.  Given that cool cores develop over a central cooling time of typically a few Gyr, one expects there to be fewer cool cores at times closer to the epoch of galaxy cluster formation.  Simulations predict a significantly higher cluster merger rate in the past \\citep{gottloeber01}.  If mergers play a strong role in the disruption of cool-core galaxy clusters, this too suggests that the fraction of galaxy clusters with cool cores should be lower at high redshifts than in local samples.  Indeed, studies of the evolution of the cool-core fraction find a much lower fraction at $z = 0.5$ \\citep{mcdonald11} and a significant dearth in {\\it strong} cool cores at $z > 0.5$ \\citep{vikhlinin07,santos10,samuele11}. To date, only a small number of $z > 0.5$ galaxy clusters with possible strong cool cores have been reported \\citep[e.g.,][]{siemiginowska10,russell12}, with the most dramatic, confirmed strong cool core coming from the South Pole Telescope (SPT) survey \\citep{mcdonald12}.  Understanding cool-core evolution is complicated by selection biases of cluster samples.  One can generically expect an X-ray-selected sample to be biased toward selecting cool-core galaxy clusters \\citep{hudson10,eckert11} due to their higher X-ray surface brightness and luminosity as compared to a noncool-core galaxy cluster of the same mass \\citep[e.g.,][]{ohara06}.  However, there are competing effects due to X-ray emission from AGNs, which are expected to be more prevalent at higher redshifts \\citep{russell12}.  The bias may be complicated further by the ways in which different cluster-finding algorithms differentiate between point sources (e.g., AGNs, X-ray binaries) and extended sources (e.g., nearby galaxies, groups, and clusters).  Furthermore, the classification of the cool-core strength of a cluster varies between surveys of different angular resolution and the method used to characterize the cool core.  Given the complex effects associated with X-ray selection, it is important to investigate the cool-core fraction and its evolution using an independent selection method.  In this paper, we study the cool-core properties of galaxy clusters selected from their Sunyaev-Zel'dovich (SZ) effect \\citep{sunyaev72} signature in the SPT survey.  At $z > 0.3$, this SPT selection is nearly redshift-independent and nearly constant in mass \\citep[e.g.,][]{reichardt12c}. The SZ effect selection is expected to be relatively insensitive to non-gravitational physics \\citep{nagai06}, the dynamical state of clusters \\citep{jeltema08}, radio contamination from point sources and BCGs \\citep{lin09}, and the presence of cool cores \\citep{Motl05}.  Therefore, an SZ-cluster survey is expected to be a useful tool to study the redshift evolution of galaxy cluster properties.  This work provides the first glimpse of the cool-core properties of a sample of galaxy clusters selected from the SZ effect.  The layout of this paper is as follows.  In Section~\\ref{sec:observations}, we detail the observations used and describe the data reduction procedures.  In Section~\\ref{sec:character-cc}, we present the steps used to make our measurements as well as compare two methods used to characterize cool-core strengths that are suitable for high-redshift clusters.  In Section~\\ref{sec:cc-evolution}, we present the results of our measurements and investigate the implications for the cool-core fraction at high redshifts.  In Section~\\ref{sec:bcg-offset}, we investigate the relationship between a cluster's cool-core strength and the offset of its BCG.  Finally, in Section~\\ref{sec:conclusions}, we conclude our analyses and present future studies and applications.  In this analysis, we assume the best-fit WMAP7+BAO+$H_0$ flat $\\Lambda$CDM cosmology \\citep{komatsu11} with Hubble parameter $H_0 = 70.4$~km~s$^{-1}$~Mpc$^{-1}$, matter density $\\Omega_M = 0.272$, and dark energy density $\\Omega_\\Lambda = 0.728$. ",
        "conclusions": ""
    },
    "1208/1208.4581_arXiv.txt": {
        "abstract": "We present results from a study of a nuclear emission of a nearby radio galaxy, 4C+29.30, over a broad 0.5--200\\,keV X-ray band. This study used new {\\it XMM-Newton} ($\\sim 17$\\,ksec) and {\\it Chandra} ($\\sim 300$\\,ksec) data, and archival {\\it Swift}/BAT data from the 58-month catalog. The hard ($>2$\\,keV) X-ray spectrum of 4C+29.30 can be decomposed into an intrinsic hard power-law ($\\Gamma\\sim1.56$) modified by a cold absorber with an intrinsic column density $N_{\\rm H,\\,z}\\sim5\\times10^{23}$\\,cm$^{-2}$, and its reflection ($|\\Omega/2\\pi| \\sim 0.3$) from a neutral matter including a narrow iron K$\\alpha$ emission line at the rest frame energy $\\sim 6.4$\\,keV. The reflected component is less absorbed than the intrinsic one with an upper limit on the absorbing column of $N^{\\rm refl}_{\\rm H,\\,z}<2.5\\times 10^{22}$\\,cm$^{-2}$. The X-ray spectrum varied between the {\\it XMM-Newton} and {\\it Chandra} observations. We show that a scenario invoking variations of the normalization of the power-law is favored over a model with variable intrinsic column density. X-rays in the 0.5--2\\,keV band are dominated by diffuse emission modeled with a thermal bremsstrahlung component with temperature $\\sim 0.7$\\,keV, and contain only a marginal contribution from the scattered power-law component. We hypothesize that 4C+29.30 belongs to a class of `hidden' AGN containing a geometrically thick torus. However, unlike the majority of them, 4C+29.30 is radio-loud. Correlations between the scattering fraction and Eddington luminosity ratio, and the one between black hole mass and stellar velocity dispersion, imply that 4C+29.30 hosts a black hole with $\\sim 10^8$\\,M$_{\\odot}$ mass. ",
        "introduction": "4C+29.30 is a low-redshift ($z=0.0647$) radio source with a moderate radio luminosity $\\sim 10^{42}$\\,erg\\,s$^{-1}$ (Siemiginowska et al. 2012) hosted by an elliptical galaxy. It was first studied in the radio and optical bands by van Breugel et al. (1986). Observed complex radio morphology was resolved into jets, knots, lobes, and a diffuse tail. A number of optical emission lines were measured, and their relative intensities indicated photoionization by a non-stellar continuum attributed to the central active galactic nucleus (AGN). The radio emission and the optical line emitting gas were found to be strongly linked suggesting interaction between the radio source and its environment. In addition, shells and dust lanes observed in the optical band (van Breugel et al. 1986) provide evidence of a past merger with a gas-rich disk galaxy which possibly triggered the rejuvenated AGN activity. Subsequently, follow-up radio studies revealed radio structures with diverse spectral ages suggesting an intermittent nature of the radio source in 4C+29.30. Jamrozy et al. (2007) presented evidence for a large scale extended relic radio emission most probably due to an earlier cycle of activity of the source $\\gtrsim 200$\\,Myr ago. The age of a small-scale radio structure embedded in the extended relic radio emission was estimated at $\\lesssim 100$\\,Myr, with the inner double knots of a spectral age of $\\lesssim 33$\\,Myr (Jamrozy et al. 2007). Liuzzo et al. (2009) resolved the very central region of the source into two nuclear knots with spectral ages of $\\sim 15$\\,yr and $\\sim 70$\\,yr. 4C+29.30 was cataloged as an infrared point-source by the IRAS, WISE and 2MASS surveys (NASA/IPAC Infrared Science Archive).  The first detection of the source in the X-ray band occurred during a snap-shot 8\\,ksec {\\it Chandra} observation (Gambill et al. 2003). The short exposure time and few detected counts prevented detailed spectral analysis. Nevertheless, this pilot observation hinted towards a complex X-ray morphology composed of a nucleus, hotspots, a putative jet, and diffuse emission. The source was not a target of X-ray pointings prior to the first {\\it Chandra} observation, without detection in, e.g., the ROSAT All Sky Survey, probably due to significant intrinsic absorption (Gambill et al. 2003, Siemiginowska et al. 2012). In 2008, we observed 4C+29.30 with {\\it XMM-Newton} for $\\sim 45$\\,ksec. This observation was followed by our deep 300\\,ksec {\\it Chandra} exposure (Siemiginowska et al. 2012, hereafter Paper I) which clearly revealed diverse X-ray emitting structures, many of which correspond to those in radio or optical bands, and others are intrinsic to the X-ray band. The soft 0.5--2\\,keV X-ray image showed northern and southern lobes, hotspots, southern jet, and thermal diffuse emission in the center. Emission in the hard X-ray band (above 2\\,keV) was instead dominated by the active nucleus of 4C+29.30. In Paper I, our discussion focused on the spectral properties of the soft X-ray components, the structure of the outflow, and interactions of the X-ray jet with the interstellar medium. We emphasized similarities between 4C+29.30 and famous radio galaxies showing jet-ambient medium interactions, i.e. M87 (e.g., Million et al. 2010), NGC 1275 (e.g., Fabian et al. 2011), and Cen A (e.g., Morganti et al. 2010, and references therein).  In the present paper, we focus on detailed modeling of the X-ray core emission to uncover the nature of the AGN powering 4C+29.30. We analyze new {\\it Chandra} data and archival {\\it Swift}/BAT data from the 58-months catalog (Baumgartner et al. 2010), and thus we study the core emission over a broad 0.5--200\\,keV X-ray band. We apply our best fit model to our new {\\it XMM-Newton} observations of 4C+29.30 taken $\\sim 2$ years before our {\\it Chandra} observation. We search for signatures of the X-ray spectral variability between the {\\it Chandra} and {\\it XMM-Newton} observations. The paper is organized as follows. In Section 2 we describe our data and provide details on the preparation of the spectra for modeling. In Section 3 we test various spectral scenarios for AGN X-ray activity. We compare the best fitting {\\it Chandra} and {\\it XMM-Newton} models in order to study the origins of possible spectral variability. In Section 4 we present the discussion of our results indicating that 4C+29.30 is a strongly absorbed intrinsically X-ray variable AGN with Seyfert 2 type of activity and a peculiar obscuring geometry, probably due to a geometrically thick torus similar to that reported for several {\\it Swift}/BAT selected AGN (e.g., Ueda et al. 2007; Winter et al. 2009a,b). Finally, in Section 5 we summarize and conclude our study. Throughout the paper we adopt $\\Omega_\\Lambda=0.73$, $\\Omega_{\\rm M}=0.27$ and $H_{\\rm 0}=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ for the flux-luminosity transformations. ",
        "conclusions": "We confirmed that the X-ray emission of the AGN in 4C+29.30, a radio source hosted by an elliptical galaxy, resembles Seyfert 2 type of activity. We estimated that the mass of the central black hole is of the order of 10$^8$\\,M$_{\\odot}$. We decomposed the broad band 0.5--200\\,keV X-ray continuum ({\\it Chandra}/{\\it XMM-Newton} supplemented by {\\it Swift}/BAT observations) into (i) a hard ($\\Gamma = 1.5$--1.6) cut-off power-law component attenuated due to a strong intrinsic X-ray absorption with column density $N_{\\rm H} \\sim 5 \\times 10^{23}$\\,cm$^{-2}$; (ii) a reflection hump ($|\\Omega/2\\pi|\\sim0.3$) and narrow iron K$\\alpha$ line detected for the first time in this source originating due to the reprocessing of the AGN power-law emission by a distant neutral matter; (iii) a soft thermal emission of a diffuse gas clearly detected in our {\\it Chandra} image dominating the 0.5--2\\,keV band; (iv) a marginal contribution of a scattered power-law emission to the soft X-ray band, $f_{\\rm sc} < 0.25$\\%. We hypothesized that the distant reflector and the obscuring matter can be both associated with a molecular torus. Based on $f_{\\rm sc}$, the ratio of the observed soft-to-hard X-ray fluxes, and the {\\it XMM-Newton} colors we proposed that 4C+29.30 belongs to the hidden/buried AGN, a new class that is emerging among the {\\it Swift}/BAT hard X-ray selected AGN. We demonstrated that 4C+29.30, a radio-loud source with an intermittent jet activity and $f_{\\rm sc}$ below 0.5\\%, shows properties characteristic to only $\\sim 9$\\% of the presently identified hidden/buried AGN candidates with estimated radio loudness parameter.\\\\"
    },
    "1208/1208.1471_arXiv.txt": {
        "abstract": "Rich young stellar clusters produce H\\,{\\sc{ii}} regions whose expansion into the nearby molecular cloud is thought to trigger the formation of new stars. However, the importance of this mode of star formation is uncertain. This investigation seeks to quantify triggered star formation (TSF) in IC 1396A (a.k.a., the Elephant Trunk Nebula), a bright rimmed cloud (BRC) on the periphery of the nearby giant H\\,{\\sc{ii}} region IC 1396 produced by the Trumpler 37 cluster. X-ray selection of young stars from {\\it Chandra} X-ray Observatory data is combined with existing optical and infrared surveys to give a more complete census of the TSF population. Over 250 young stars in and around IC 1396A are identified; this doubles the previously known population.  A spatio-temporal gradient of stars from the IC 1396A cloud toward the primary ionizing star HD 206267 is found.  We argue that the TSF mechanism in IC 1396A is the radiation-driven implosion process persisting over several million years. Analysis of the X-ray luminosity and initial mass functions indicates that $>140$ stars down to 0.1~M$_{\\odot}$ were formed by TSF. Considering other BRCs in the IC 1396 H\\,{\\sc{ii}} region, we estimate the TSF contribution for the entire H\\,{\\sc{ii}} region exceeds $14-25$\\% today, and may be higher over the lifetime of the H\\,{\\sc{ii}} region. Such triggering on the periphery of H\\,{\\sc{ii}} regions may be a significant mode of star formation in the Galaxy. ",
        "introduction": "\\subsection{Motivation}\\label{motivation_subsection} A long-standing issue in Galactic star formation is: How much star formation relies on spontaneous gravitational collapse of cold clouds? $vs.$ How much star formation relies on collapse ``triggered'' by the imposition of external forces? Triggered star formation (TSF) may play roles on large-scales from compression by Galactic spiral arms, on meso-scales from compression by supernova remnant (SNR) superbubbles in starburst complexes, and on small-scales from compression by expanding H\\,{\\sc{ii}} regions \\citep{Elmegreen07}. It has long been recognized that star formation in molecular clouds can be triggered by ionization or shock fronts produced by nearby massive stars \\citep{Elmegreen77}. Two major mechanisms for triggering by expanding H\\,{\\sc{ii}} regions have been widely discussed: the Radiation Driven Implosion (RDI) model and the Collect and Collapse (C\\&C) model.  In the RDI model, the expanding OB ionization front ablates the surface of surrounding pre-existing cloudlets producing bright-rimmed clouds (BRCs), often with cometary structures, and driving inward a compressional shock that induces star formation. Discussed since the 1980s, RDI now has well-developed hydrodynamical calculations showing the viability of TSF on timescales $\\sim 0.1$~Myr \\citep[][and references therein]{Kessel03,Miao09,Bisbas11}.  Observationally, two types of bright-rimmed clouds (BRCs) can be distinguished: relatively massive cloudlets with widths of $w>0.5-0.7$~pc, which are often found to harbor more than a few stars at their tips; and thin ($w<0.5$~pc) elongated pillar-like structures with a few or no stars at their tips \\citep{Chauhan11}. According to the RDI model, this distinction may reflect different evolutionary stages of BRCs with narrow pillars interpreted as the last vestiges of the photoevaporated and dispersed cloudlet that has insufficient amount of remaining gas to form new stars \\citep{Bisbas11}. Alternative explanations for the thin elongated pillars involve hydrodynamic instabilities in ionization fronts \\citep{Mizuta06,Whalen08} or ionization of large turbulent clouds without pre-existing cloudlets \\citep[][and references therein]{Gritschneder10}.  In most cases, molecular, infrared (IR), and H$\\alpha$ surveys of BRCs trace only the most recently formed stars \\citep[e.g.,][]{Sugitani95,Ogura02,Urquhart09}. While longer wavelength surveys select principally the youngest stars undergoing accretion from dusty disks, X-ray surveys efficiently select young stars due to enhanced magnetic reconnection flaring that extends to the main sequence \\citep{Feigelson10}. In a few cases, $Chandra$ observations have added the diskfree pre-main sequence (PMS) populations to TSF clouds \\citep[e.g.,][]{Getman07, Getman09}. Triggered BRCs often show embedded mid-IR sources denoting protostars, while H$\\alpha$, $JHK$, and X-ray surveys reveal small clusters of disky and diskfree PMS stars within and in front of the bright rim\\footnote{Henceforth, the term ``in front of'' refers to the region between a globule and an ionizing star(s) in 2-dimensional projection, and not the region between the observer and the globule.}. In a few cases, spatial-age gradients in the stellar population are clearly seen where the youngest stars are embedded and older stars are aligned toward the ionizing sources \\citep[e.g.,][]{Ogura07, Getman07, Ikeda08, Getman09}.  This directly supports the RDI mechanism and implies that the existing clouds have been actively forming stars for more than 1~Myr.  In the C\\&C model, ionization and wind shock fronts sweep up neutral material into a dense shell that becomes gravitationally unstable and fragments. Individual stars or entire massive clusters may form inside these fragments. The theory is moderately developed \\citep{Whitworth94, Hosokawa06, Dale07} and may explain a number of spectacular examples of embedded clusters found around H\\,{\\sc{ii}} regions using long-wavelength (mm to near-IR) observations \\citep{Deharveng03, Zavagno06, Zavagno07}. They find a central ionizing star with a roughly spherical H\\,{\\sc{ii}} region surrounded by a shell of dense molecular gas with dense fragments harboring young stellar objects. However, recent SPH simulations by \\citet{Walch11} show that if an H\\,{\\sc{ii}} region expands into a clumpy molecular cloud then the production of a molecular shell with dense massive fragments does not require the C\\&C mechanism. Instead, triggered star formation would take place through simultaneous enhancement of density and global radiation-driven implosion of pre-existing molecular clumps.  Star formation triggered by H\\,{\\sc{ii}} regions may be very common. The {\\it Spitzer}/GLIMPSE survey finds $\\sim$10 mid-IR bubbles/deg$^{2}$ throughout the Galactic Disk \\citep{Churchwell06, Simpson12}. \\citet{Deharveng10} show that most of these bubbles enclose H\\,{\\sc{ii}} regions ionized by O-B2 stars, many are surrounded by cold dust shells and many have bright-rimmed dust condensations protruding inside the H\\,{\\sc{ii}} region. $26$\\% of the dust shell condensations have ultracompact H\\,{\\sc{ii}} (UCH\\,{\\sc{ii}}s) regions and/or methanol masers, likely indicating triggering of massive star formation. Based on the cross-correlation of the positions of the mid-IR shells with the positions of the very bright and red young stellar objects (YSOs) from the Red MSX Source (RMS) catalog, \\citet{Thompson11} estimate that the formation of $14-30$\\% of massive stars in the Milky Way could be triggered by expanding H\\,{\\sc{ii}} regions.  \\begin{figure*} \\centering \\includegraphics[angle=0.,width=170mm]{fig1.eps} \\caption{(a) A $3\\degr \\times 3\\degr$ image of the emission nebula IC 1396 from the Digitized Sky Survey (DSS). The primary ionizing source of the region, the O6 star HD 206267, and the bright-rimmed globule IC 1396 are marked in red. The blue polygon outlines the ACIS field centered on the IC 1396A. (b) Close-up, combined X-ray and IR image of IC 1396A. The adaptively smoothed {\\it Chandra}-ACIS image in the $0.5-8.0$~keV band (blue) is superposed on the $Spitzer$-IRAC composite image in the 3.6~$\\mu$m (green) and 8.0~$\\mu$m (red) bands. The  ACIS field is outlined in blue; the head and the trunk of the cloud are marked in red. The positions of the two brightest features of the ionized optical rim in front of the cloud are marked in red as Rims A and Aa. North is up, east is to the left \\label{fig_intro}} \\end{figure*}  Due to violent energy feedback from massive stars in large star-forming complexes and individual H\\,{\\sc{ii}} regions, the geometry, conditions, and history of star formation are often difficult to ascertain. The principal signatures of small-scale TSF --- a star cluster in and around the cloudlet on the periphery of an expanding HII region --- may be erased through removal of the cloudlet by stellar winds and UV radiation, or the kinematic dispersal of the unbound cluster within less than a few million years. Thus, it is often difficult to identify sites of triggered star formation and to quantify the impact of triggering processes. IC 1396 is a nearby \\citep[$D \\sim 870$~pc;][]{Contreras02}, large ($\\sim 12$~pc in radius) shell-like H\\,{\\sc{ii}} region, where traces of recent triggered star formation are still evident \\citep[][and references therein]{Kun08}. The goal of the current study is to identify, understand, and quantify the triggered stellar population in the central part of the region associated with the largest BRC, IC 1396A (The Elephant Trunk Nebula), and to estimate the contribution of triggered star formation to the total stellar population in both the central part and, with some additional assumptions, in the entire H\\,{\\sc{ii}} region.   \\subsection{The target: IC~1396A globule}\\label{ic1396_subsection}  The IC 1396 H\\,{\\sc{ii}} region, Sh 2-131 at $(l,b)=(99.3^\\circ,3.7^\\circ)$, is excited mainly by the O6.5f star HD 206267 located at the center of the region in the 4~Myr old cluster Tr~37 \\citep{Sicilia-Aguilar05}[hereafter SA05]. IC 1396 has a rich population of $>20$ bright-rimmed and cometary globules seen in silhouette against the emission nebula \\citep[][and Figure \\ref{fig_intro} here]{Sugitani91, Froebrich05}. Many of the clouds reside on the large molecular shell surrounding the H\\,{\\sc{ii}} region, which expands at a speed of $5$~km~s$^{-1}$ with an inferred expansion time around $2.5$~Myr \\citep{Patel95}. Sites of possible star-formation have been identified in/around at least several globules \\citep[e.g.][]{Schwartz91, Ogura02, Froebrich05} including sites of substantial star-formation in IC~1396N \\citep{Nisini01, Getman07, Beltran09, Choudhury10}, SFO~37 \\citep{Ikeda08}, and IC~1396A \\citep{Reach04, Sicilia-Aguilar05, Sicilia-Aguilar06, Morales-Calderon09, Barentsen11}.  Lying $\\sim 15\\arcmin$ ($\\sim 3.7$~pc projected distance) west of HD~206267, IC 1396A is the bright-rimmed cloud closest on the sky to HD~206267. Its optical rim is the brightest of all rims and the cloud is thus likely to be the physically closest to HD 206267 \\citep{Weikard96}. Following \\citet{Weikard96} we designate the two brightest features of the rim as Rim A located in front of the $>1$~pc diameter head of the globule and Rim Aa located further west at the shoulder (trunk) of the cloud (Figure \\ref{fig_intro}). The part of the cloud behind Rim Aa has its own designation: SFO~36 \\citep{Sugitani91}. A 0.3~pc diameter cavity in the molecular cloud is present in the head of the globule produced by the Herbig Ae star LkHa~349 \\citep{Hessman95}. Outside this ``hole'', the globule is optically thick ($A_V \\sim 10$~mag) with molecular mass $\\sim 200$~M$_{\\odot}$ \\citep{Patel95, Weikard96}.  Early {\\it Spitzer} images revealed both bright diffuse mid-infrared (MIR) emission and 17 Class~II/I stars in the IC 1396A globule \\citep[][their Figure 2]{Reach04}.  A spectacular multi-band movie of the cloud was one of the first press releases of the {\\it Spitzer} mission\\footnote{The press release can be found at http://www.spitzer.caltech.edu/images/1058-ssc2003-06b-Dark-Globule-in-IC-1396}. SA05 and \\citet{Sicilia-Aguilar06} combined new deeper and wider-area {\\it Spitzer} photometry with optical photometry and spectroscopy to reveal a rich lightly obscured population of $>200$ H$\\alpha$ emission and/or Li~6707~\\AA\\ absorption stars. The population is a mixture of $\\sim 4$~Myr old members of the central Tr~37 cluster and $\\sim 1$~Myr stars spatially concentrated in an arc in front of the globule (see Figure 11 in SA05). In addition, $\\sim 50$ heavily obscured IR-excess-selected objects were found inside the cloud \\citep{Sicilia-Aguilar06}. Deeper {\\it Spitzer}-IRAC maps from the time-series monitoring of the globule \\citep{Morales-Calderon09} [hereafter MC09] uncovered an additional dozen heavily embedded Class I/II YSOs. The spatial distribution of sixty five disky members in/around the globule is shown in Figure 16 of MC09. Recently, using the data from the INT Photometric H-Alpha Survey (IPHAS), \\citet{Barentsen11} [hereafter B11] identified about two dozen low-mass high-accretion T-Tauri systems spatially concentrated in an arc in front of the globule (see Figure 18a in B11). About half of these were newly discovered stellar members of the region. Spatial clustering of optical stars with younger ages and increasing accretion rates away from the ionizing star HD 206267 (SA05, B11) led SA05 and B11 to propose the presence of a stellar population in front of the IC 1396A globule whose formation has been triggered by HD 206267. Thus, several high-sensitivity optical and mid-infrared surveys have revealed a few hundred young stars in the Tr 37 region, with the youngest concentrated within and around IC 1396A.  \\begin{table*}\\scriptsize \\centering \\begin{minipage}{180mm} \\caption{Basic X-ray Source Properties. This table is available in its entirety in the electronic edition of the journal. A portion is shown here for guidance regarding its form and content. Source net counts, background counts, median energy, and apparent photometric flux are given for the full $(0.5-8)$~keV band. Column 1: X-ray source number. Column 2: IAU designation. Columns 3-4: Right ascension and declination for epoch J2000.0 in degrees. Column 5: 1~$\\sigma$ error circle around source position. Column 6: Average off-axis angle for the two merged X-ray observations. Column 7: Source net counts in merged apertures and their 1~$\\sigma$ upper errors. Column 8: Observed background counts in merged apertures. Column 9: Average PSF fraction for merged observations (at 1.5~keV) enclosed within source aperture. Column 10: Smallest of the three $p$-values (for the full, soft, and hard energy bands) for no-source hypothesis. Column 11: Effective exposure time. Columns 12: Smallest $p$-value for the one-sample Kolmogorov-Smirnov statistic under the no-variability null hypothesis within a single-observation. Column 13: $p$-value for the one-sample Kolmogorov-Smirnov statistic under the no-variability null hypothesis over the combined observations. Columns 14-15: X-ray median energy and apparent photometric flux.} \\label{tbl_xray_photometry} \\begin{tabular}{@{\\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }c@{ \\vline }} \\cline{1-15} &&&&&&&&&&&&&&\\\\ No. & CXOIC1396A & R.A. & Decl. & PosErr & $\\Theta$ & $NC$ & $BC$ & {PSF} & PbNoSrc & EffExp & PbVar1 & PbVar2 & $ME$ & $F_X$\\\\ &&(deg)&(deg)&($\\arcsec$)&($\\arcmin$)&(cnts)&(cnts)&&&(ks)&&&(keV)&(ph/cm$^2$/s)\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)&(14)&(15)\\\\ \\cline{1-15} &&&&&&&&&&&&&&\\\\ 1 & 213529.03+573412.3 & 323.871000 &  57.570087 & 0.5 & 11.1 &  $  130.9 \\pm 12.9 $  & 10.1 & 0.90 &  0.0E+00 & 28.0 &  1.8E-01 &  1.8E-01 & 2.3 &  2.6E-05\\\\ 2 & 213532.88+572846.9 & 323.887026 &  57.479703 & 0.3 & 10.8 &  $  271.5 \\pm 17.9 $  & 11.5 & 0.91 &  0.0E+00 & 29.1 &  3.2E-01 &  3.2E-01 & 1.5 &  4.9E-05\\\\ 3 & 213537.29+572842.0 & 323.905409 &  57.478345 & 1.2 & 10.2 &  $   16.3 \\pm  6.3 $  & 10.7 & 0.91 &  6.7E-07 & 29.1 &  6.9E-01 &  6.9E-01 & 1.4 &  2.9E-06\\\\ 4 & 213538.86+573032.7 & 323.911957 &  57.509103 & 0.9 &  9.8 &  $   27.7 \\pm  7.2 $  &  9.3 & 0.90 &  1.3E-12 & 28.9 &  8.9E-01 &  8.9E-01 & 2.0 &  5.0E-06\\\\ 5 & 213539.72+573127.5 & 323.915509 &  57.524320 & 1.1 &  9.1 &  $   13.7 \\pm  5.6 $  &  6.3 & 0.89 &  6.2E-07 & 28.8 &  6.2E-01 &  6.2E-01 & 3.8 &  2.5E-06\\\\ 6 & 213541.87+573252.8 & 323.924497 &  57.548001 & 1.2 &  9.1 &  $    9.3 \\pm  5.0 $  &  5.7 & 0.89 &  1.6E-04 & 28.9 &  9.9E-01 &  9.9E-01 & 2.9 &  1.7E-06\\\\ 7 & 213542.94+573334.8 & 323.928939 &  57.559685 & 0.9 &  9.1 &  $   19.1 \\pm  6.2 $  &  6.9 & 0.90 &  4.9E-12 & 28.9 &  6.0E-01 &  6.0E-01 & 1.1 &  3.4E-06\\\\ 8 & 213543.96+572955.9 & 323.933171 &  57.498871 & 0.8 &  8.9 &  $   21.4 \\pm  6.8 $  & 10.6 & 0.90 &  1.4E-08 & 48.4 &  2.0E-01 &  2.4E-01 & 1.7 &  2.2E-06\\\\ 9 & 213545.27+573326.5 & 323.938631 &  57.557375 & 1.0 &  8.8 &  $   11.9 \\pm  5.4 $  &  6.1 & 0.90 &  4.2E-07 & 27.7 &  7.7E-01 &  7.7E-01 & 1.4 &  2.2E-06\\\\ 10 & 213550.45+573547.5 & 323.960213 &  57.596543 & 0.4 &  9.2 &  $  110.5 \\pm 11.8 $  &  4.5 & 0.90 &  0.0E+00 & 20.7 &  2.3E-01 &  2.3E-01 & 1.5 &  2.7E-05\\\\ &&&&&&&&&&&&&&\\\\ \\cline{1-15} \\end{tabular} \\end{minipage} \\end{table*}  \\begin{table*}\\tiny \\centering \\begin{minipage}{180mm} \\caption{Optical and IR Photometry of X-ray Sources. This table is available in its entirety in the electronic edition of the journal. A portion is shown here for guidance regarding its form and content. Column 1: X-ray source number.  Columns 2-5: Optical $V_{FLWO}$, $I_{C,FLOW}$, $V_{LAICA}$, and $I_{J,LAICA}$  magnitudes. Columns 6-9: $2MASS$ $JHK_s$ magnitudes, and $2MASS$ photometry quality and confusion-contamination flags. Columns 10-14: IRAC magnitudes, and a digital flag giving photometric apertures and level of source contamination from nearby sources and nebular IR emission: ``3'' and ``4'' --- 3-pixel (2.6$\\arcsec$) and 4-pixel (3.46$\\arcsec$) apertures with contaminating flux from a neighboring source of no more than $5$\\%; ``2'' and ``2c'' --- 2-pixel (1.73$\\arcsec$) apertures with possible contaminating flux of $<10$\\% and $>10$\\%, respectively.} \\label{tbl_oir_photometry} \\begin{tabular}{@{\\vline}|c@{ \\vline}|c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}} \\cline{1-14} &&&&&&&&&&&&&\\\\ No. & $V_{FLWO}$ & $I_{C,FLWO}$ & $V_{LAICA}$ & $I_{J,LAICA}$ & $J$ & $H$ & $K_s$ & F$_2$ & [3.6] & [4.5] & [5.8] & [8.0] & F$_3$\\\\ &(mag)&(mag)&(mag)&(mag)&(mag)&(mag)&(mag)&&(mag)&(mag)&(mag)&(mag)&\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)&(14)\\\\ &&&&&&&&&&&&&\\\\ \\cline{1-14} &&&&&&&&&&&&&\\\\ 1 & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ...\\\\ 2 & ... & ... & ... & ... &  $  12.64 \\pm   0.03 $  &  $  11.82 \\pm   0.03 $  &  $  11.56 \\pm   0.02 $  & AAA000 & ... & ... & ... & ... & ...\\\\ 3 & ... & ... & ... & ... &  $  13.65 \\pm   0.03 $  &  $  12.95 \\pm   0.03 $  &  $  12.71 \\pm   0.03 $  & AAA000 & ... & ... & ... & ... & ...\\\\ 4 & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ...\\\\ 5 & ... & ... & ... & ... & ... & ... & ... & ... &  $  16.20 \\pm   0.06 $  &  $  15.64 \\pm   0.05 $  & ... & ... & 2\\\\ 6 &  $13.44 \\pm  0.00 $  &  $12.37 \\pm  0.00 $  &  $13.57 \\pm  0.00 $  &  $12.01 \\pm  0.00 $  &  $  11.49 \\pm   0.03 $  &  $  11.18 \\pm   0.03 $  &  $  11.07 \\pm   0.03 $  & AAA000 &  $  10.92 \\pm   0.01 $  &  $  10.90 \\pm   0.02 $  &  $  10.92 \\pm   0.04 $  &  $  10.75 \\pm   0.03 $  & 3\\\\ 7 & ... & ... &  $19.48 \\pm  0.01 $  &  $15.51 \\pm  0.00 $  &  $  14.30 \\pm   0.04 $  &  $  13.44 \\pm   0.03 $  &  $  13.07 \\pm   0.04 $  & AAA000 &  $  12.53 \\pm   0.00 $  &  $  12.28 \\pm   0.00 $  &  $  12.13 \\pm   0.01 $  &  $  11.54 \\pm   0.04 $  & 4\\\\ 8 & ... & ... &  $19.44 \\pm  0.01 $  &  $16.48 \\pm  0.00 $  &  $  14.22 \\pm   0.04 $  &  $  13.38 \\pm   0.04 $  &  $  13.17 \\pm   0.04 $  & AAA000 &  $  12.86 \\pm   0.05 $  &  $  12.82 \\pm   0.06 $  &  $  12.65 \\pm   0.05 $  & ... & 2\\\\ 9 &  $19.50 \\pm  0.08 $  &  $16.41 \\pm  0.01 $  &  $19.36 \\pm  0.01 $  &  $15.71 \\pm  0.00 $  &  $  14.61 \\pm   0.04 $  &  $  13.87 \\pm   0.04 $  &  $  13.66 \\pm   0.04 $  & AAA000 &  $  13.34 \\pm   0.01 $  &  $  13.24 \\pm   0.01 $  &  $  13.26 \\pm   0.07 $  & ... & 4\\\\ 10 & ... & ... & ... & ... &  $  12.86 \\pm   0.03 $  &  $  12.28 \\pm   0.04 $  &  $  12.11 \\pm   0.03 $  & AAAcc0 & ... & ... & ... & ... & ...\\\\ &&&&&&&&&&&&&\\\\ \\cline{1-14} \\end{tabular} \\end{minipage} \\end{table*}  As X-ray emission from PMS stars is based on enhanced solar-type magnetic reconnection events rather than disk or accretion processes, X-ray selection delivers rich and clean samples of diskless stars missed by H$\\alpha$ and IR selection \\citep{Feigelson10}. Using archived X-ray {\\it Chandra} grating data, \\citet{Mercer09} identifed 22 new PMS stars around HD 206267, more than doubling the number of previously reported young stars in the central $10\\arcmin \\times 8\\arcmin$ area of Tr~37. We present here the identification of $>130$ previously unknown members of the Tr 37/IC 1396A stellar populations, using new X-ray and optical observations with archived {\\it Spitzer} data. The {\\it Chandra} and the auxiliary {\\it Spitzer} and optical FLWO/LAICA data are described in \\S \\ref{data_reduction_section}. Identification of X-ray sources with optical-IR counterparts and membership classification are considered in \\S \\ref{ir_optical_counterparts_section}. Classification of diskbearing and diskless X-ray emitting YSOs is provided in \\S \\ref{disk_classes_subsection}, and the discovery of a new population of non-{\\it Chandra} IR-excess members is given in \\S \\ref{non_chandra_members_subsection}. We then estimate the global properties of the Tr~37/IC~1396A stellar populations including age distribution (\\S \\ref{age_analysis_section}), spatial structure (\\S \\ref{spatial_structure_section}), X-ray luminosity functions (XLFs) and initial mass functions (IMFs; \\S \\ref{total_populations_section}). We end with a discussion of the implications for understanding and quantifying triggered star formation in the region (\\S \\ref{implications_for_tsf_section}).  \\begin{table*}\\tiny \\centering \\begin{minipage}{180mm} \\caption{Derived Properties, Stellar Identifications, Membership, and Classification. This table is available in its entirety in the electronic edition of the journal. A portion is shown here for guidance regarding its form and content. Column 1: X-ray source number. Columns 2-3: X-ray net counts and median energy for the full $(0.5-8.0)$~keV band. Columns 4-5: X-ray column density, intrinsic luminosity, and their errors (summed in quadrature statistical and systematic errors). Luminosities are derived assuming a distance of $870$~pc and are given for the full $(0.5-8.0)$~keV band. Columns 6-7: Age estimates derived from the optical $V$ $vs.$ $V-I_C$ and $V$ $vs.$ $V-I_J$  color magnitude diagrams using PMS models of \\citet{Siess00} and assuming a constant source extinction for optical PMS candidates of $A_V = 1.56$~mag (SA05, B11). Column 8: Apparent SED slope from IRAC photometry with $1\\sigma$ error. Column 9: Number of IRAC bands, from which the SED slope was derived. Column 10: $2MASS$ identifier. Column 11: Source identifier from our analysis of the IRAC data; X-ray sources outside the field of view of the IRAC data analyzed here are labeled as ``OutOfIracFOV''. Columns 12-14: Stellar counterparts from SA05, B11, and MC09. Column 15: This positional flag indicates if an X-ray source lies projected against the globule (``1'') or not (``0''). Column 16: Indicates membership and YSO class: ``EXG'' --- possible extragalactic contaminant; ``FRG'' --- possible foreground contaminant or YSO; ``DSK'' and ``NOD'' --- disky and diskless YSOs, respectively; ``UNC1'' and ``UNC2'' --- objects of the uncertain class.  ``UNC1'' are sources with no IR counterparts; these are most likely extragalactic objects but some can be background stars. ``UNC2'' have weak registered and unregistered IR counterparts or lie close to bright IR objects; these could be YSOs or field stars or extragalactic objects.} \\label{tbl_derived_props} \\begin{tabular}{@{\\vline}|c@{ \\vline}|c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}c@{ \\vline}} \\cline{1-16} &&&&&&&&&&&&&&&\\\\  No. & $NC$ & $ME$ & $\\log(N_H)$ & $\\log(L_X)$ & $t_{FLWO}$ & $t_{LAICA}$ & $\\alpha_0$ & $Np$ & 2MASS & IRAC & SA05 & B11 & MC09 & Region & Class\\\\  &(cnt)&(keV)&(cm$^{-2}$)&(erg/s)&(Myr)&(Myr)&&&&&&&&&\\\\ (1)&(2)&(3)&(4)&(5)&(6)&(7)&(8)&(9)&(10)&(11)&(12)&(13)&(14)&(15)&(16)\\\\ &&&&&&&&&&&&&&&\\\\ \\cline{1-16} &&&&&&&&&&&&&&&\\\\ 164 &    7.8 & 1.8 & ... & ... & ... & ... &  $  -1.36 \\pm   0.10 $  & 4 & 21364762+5729540 & J213647.61+572954.1 & ... &  30 & IC1396A-61 & 1 & DSK\\\\ 165 &    3.4 & 2.7 & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & 0 & UNC1\\\\ 166 &   18.9 & 4.1 & $ 23.04 \\pm 0.21 $ & $ 30.97 \\pm 0.24 $ & ... & ... &  $  -1.15 \\pm   0.20 $  & 4 & 21364788+5731306 & J213647.87+573130.6 & ... & ... & IC1396A-28 & 1 & DSK\\\\ 167 &   17.9 & 1.3 & $ 21.60 \\pm 0.50 $ & $ 29.71 \\pm 0.23 $ & ... & ... &  $  -2.43 \\pm   0.07 $  & 3 & 21364793+5723062 & J213647.94+572306.5 & ... & ... & ... & 0 & NOD\\\\ 168 &   12.5 & 1.3 & ... & ... &   4.7 &   2.8 &  $  -2.54 \\pm   0.07 $  & 2 & 21364819+5734020 & J213648.19+573401.9 & ... & ... & ... & 0 & NOD\\\\ 169 &   59.0 & 1.1 & $ 20.15 \\pm 0.77 $ & $ 29.92 \\pm 0.09 $ & ... & ... &  $  -2.83 \\pm   0.01 $  & 4 & 21364825+5739185 & J213648.25+573918.4 & ... & ... & ... & 0 & FRG\\\\ 170 &   11.7 & 3.1 & $ 22.58 \\pm 0.26 $ & $ 30.61 \\pm 0.30 $ & ... & ... & ... & ... & ... & OutOfIracFOV & ... & ... & ... & 0 & UNC2\\\\ 171 &   12.8 & 1.3 & $ 21.60 \\pm 0.69 $ & $ 29.48 \\pm 0.25 $ & ... & ... &  $  -2.25 \\pm   0.11 $  & 3 & 21364866+5730262 & J213648.65+573026.2 & ... & ... & ... & 1 & NOD\\\\ 172 &    6.9 & 1.0 & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & ... & 0 & UNC1\\\\ 173 &   15.9 & 2.6 & $ 22.41 \\pm 0.17 $ & $ 30.48 \\pm 0.25 $ &   0.6 &   0.3 &  $  -0.94 \\pm   0.01 $  & 4 & 21364941+5731220 & J213649.42+573122.2 & 14-141 & ... & IC1396A-29 & 1 & DSK\\\\ 174 &   96.1 & 1.5 & $ 21.78 \\pm 0.21 $ & $ 30.56 \\pm 0.16 $ &   3.5 &   3.8 &  $  -1.76 \\pm   0.02 $  & 4 & 21364964+5722270 & J213649.65+572227.1 & ... & ... & ... & 0 & DSK\\\\ &&&&&&&&&&&&&&&\\\\ \\cline{1-16} \\end{tabular} \\end{minipage} \\end{table*} ",
        "conclusions": "Signatures of small-scale triggered star formation on the periphery of H\\,{\\sc{ii}} regions can be erased on short timescales and thus, it is often difficult to identify sites of triggered star formation and to measure the impact of triggering processes. There are only a handful of studies in the literature that have attempted to estimate the contribution of triggered star formation from expanding H\\,{\\sc{ii}} regions. Infrared emitting protostars in cloud cores and larger populations of PMS stars with spatial gradients in stellar age indicative of long-standing triggering processes have both been seen associated with several clouds on the periphery of expanding H\\,{\\sc{ii}} regions. Much of the progress is due to the fruitful combination of X-ray and optical/infrared surveys; the former captures the older diskless PMS population while the latter detects the disky and accretion population.  IC 1396 is a nearby, large shell-like H\\,{\\sc{ii}} region, where traces of recent triggered star formation are still evident. We present new {\\it Chandra} X-ray data and auxiliary IR {\\it Spitzer} and optical FLWO/LAICA data of YSOs associated with the central cluster Trumpler~37 and the adjacent bright-rimmed cloud IC 1396A. These data are merged with the data from previously published optical-IR stellar catalogs of the region. The goal of this work is to identify, understand, and quantify the triggered stellar population in the region associated with the IC 1396A cloud and to estimate the contribution of triggered star formation to the entire H\\,{\\sc{ii}} region.  Out of $415$ X-ray sources detected in the Tr~37/IC~1396A area, 175 are identified as YSOs and classified on 124 diskless and 51 disky objects (\\S \\ref{ir_optical_counterparts_section}). The majority of the remaining X-ray sources are contaminants, mainly extragalactic sources unrelated to the region (\\S \\ref{membership_subsection}). In addition to the 175 X-ray emitting YSOs, we identify 42 non-$Chandra$ IR-excess low-mass members of the region (\\S \\ref{non_chandra_members_subsection}). Combining these data with the previous optical (SA05, B11) and IR (MC09) catalogs of YSOs yields a total of 259 YSOs in/around IC~1396A cloud within the ACIS field, half of which are new members discovered in this work (\\S \\ref{spatial_structure_section}).  The age and spatial distributions of young stars reveal a spatio-temporal gradient of stars from the IC~1396A cloud toward the primary ionizing star HD 206267 (\\S \\ref{age_analysis_section} and \\ref{spatial_structure_section}). Young stars are found to be clustered along the ionized rim, both outside and inside the cloud. The clustering of young stars in front of the cloud is consistent with the previously reported arc-shaped layer of the low-mass high accretion stars (B11) and/or the $\\sim 1$~Myr old optical stars (SA05). Quantitatively, the number of stars clustered in front of the cloud is consistent with the 2-fold increase in the stellar surface density at the western border of the central Tr~37 cluster (\\S \\ref{spatial_structure_section} and \\ref{discussion_centralpart_subsection}). The stellar age increases from $\\la 1$~Myr inside the cloud, to $<2-3$~Myr in front of the cloud, to $\\sim 4$~Myr towards the central Tr~37 cluster. Combined with the results from the previous studies, these findings provide strong evidence for RDI triggered star formation over millions of years in the Tr~37/IC~1396A region (\\S \\ref{discussion_rdi_evidences_subsection}). Based on this evidence, we believe that the majority of the YSOs found inside the cloud and the majority of the stars found clustered in front of the cloud were likely formed through the RDI process extending over $2-3$~Myr.  By scaling the XLF/IMF measurements obtained for different YSO samples in the region (\\S \\ref{total_populations_section}) to the number of the observed triggered stars, we estimate a total of $>140$ triggered stars down to 0.1~M$_{\\odot}$ in/around the IC~1396A cloud, constituting $\\ga 25$\\% of the total popualtion in the central part of the region today (\\S \\ref{discussion_centralpart_subsection}). Furthermore, with a few additional assumptions we estimate the contribution of the triggered stellar population to the total population of the entire IC 1396 H\\,{\\sc{ii}} region to be $>14$\\% (\\S \\ref{discussion_entire_subsection}). Thus, the currently active clouds on the periphery of IC 1396 add at least $14-25$\\% to the current cluster population. When past and future clouds are considered, the contribution of triggering to the star formation in the molecular cloud could be much higher.  The inferred apparent value of star formation efficiency for the triggered population in IC~1396A is $30\\%$ using current gas cloud masses (\\S \\ref{discussion_sfe_subsection}). Our findings support a picture of continuing star formation over millions of years when the RDI mechanism occurs in a larger than the presently seen IC 1396A cloud composed of multiple clumps (\\S \\ref{discussion_rdi_clumps_subsection})."
    },
    "1208/1208.1192_arXiv.txt": {
        "abstract": "We construct a simple model of universe with a generalized equation of state $p=(\\alpha +k\\rho^{1/n})\\rho c^2$ having a linear component $p=\\alpha\\rho c^2$ and a polytropic component $p=k\\rho^{1+1/n}c^2$. For $\\alpha=1/3$, $n=1$ and $k=-4/(3\\rho_P)$, where $\\rho_P=5.16\\, 10^{99}\\, {\\rm g}/{\\rm m}^3$ is the Planck density, this equation of state provides a model of the early universe without singularity describing the transition between the pre-radiation era and the radiation era. The universe starts from $t=-\\infty$ but, when $t<0$, its size is less than the Planck length $l_P=1.62\\, 10^{-35}\\, {\\rm m}$. The universe undergoes an inflationary expansion that brings it to a size $a_1=2.61\\, 10^{-6}\\, {\\rm m}$ on a timescale of a few Planck times $t_P=5.39\\, 10^{-44}\\, {\\rm s}$. When $t\\gg t_P$, the universe decelerates and enters in the radiation era. For $\\alpha=0$, $n=-1$ and $k=-\\rho_{\\Lambda}$, where $\\rho_{\\Lambda}=7.02\\, 10^{-24}\\, {\\rm g}/{\\rm m}^3$ is the cosmological density, this equation of state describes the transition from a  decelerating universe dominated by baryonic and dark matter to an accelerating universe  dominated by dark energy (second inflation). The transition takes place at a size $a_2=8.95\\, 10^{25}{\\rm m}$ corresponding to a time of the order of the cosmological time $t_{\\Lambda}=1.46\\, 10^{18}\\, {\\rm s}$. The present universe turns out to be just at the transition ($t_0\\sim t_{\\Lambda})$. This polytropic model reveals a nice ``symmetry'' between the early  and late evolution of the universe, the cosmological constant $\\Lambda$ in the late universe playing a role similar to the Planck constant $\\hbar$ in the early universe. We interpret the cosmological constant as a fundamental constant of nature describing the ``cosmophysics'' just like the Planck constant describes the microphysics. The Planck density and the cosmological density represent fundamental upper and lower bounds differing by ${122}$ orders of magnitude. The cosmological constant ``problem'' may be a false problem. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.2702_arXiv.txt": {
        "abstract": "Recent high-sensitivity observation of the nearby radio galaxy M87 have provided important insights into the central engine that drives the large-scale outflows seen in radio, optical and X-rays. This review summarizes the observational status achieved in the high energy (HE;$<$ 100 GeV) and very high energy (VHE; $>$100 GeV) gamma-ray domains, and discusses the theoretical progress in understanding the physical origin of this emission and its relation to the activity of the central black hole. ",
        "introduction": "The advanced capabilities of current gamma-ray instruments offer a unique tool to elucidate the nature of the central engine in extragalactic objects. At the time of writing, about 1000 active galaxies have been detected above 100 MeV by Fermi-LAT.\\cite{ackermann11} Ground-based gamma-ray instruments on the other hand have established TeV emission and spectra for about 50 active galaxies.\\footnote{For an up-to-date list see the online TeV Source Catalog (TeVCat) at http://tevcat.uchicago.edu/} The majority of these are blazars, only a minor fraction are radio galaxies. The giant elliptical Virgo-cluster galaxy M87 (NGC 4486, 3C274) is particularly unique among those sources, being the nearest AGN in the northern sky, one of the first extragalactic objects discovered (in 1918) to have an optical jet\\cite{curtis1918}, the first extragalactic X-ray source to be identified\\cite{bradt67}, and the first non-blazar active galaxy detected at VHE energies.\\cite{aharonian03} Because of its proximity, many astrophysical phenomena can be studied in exceptional detail, which has led M87 to become a prominent benchmark for theoretical research. In particular, since the jet in M87 is substantially inclined, it is fairly easy to study its core - the vicinity of the black hole. ",
        "conclusions": "The radio galaxy M87 is a unique source on the extragalactic sky. Its proximity and huge central black hole mass allow for an analysis of fundamental astrophysical processes in unprecedented detail. M87 is well-known for a one-sided, kpc-size relativistic jet visible at radio, optical and X-ray wavelengths. However, compared to typical blazar sources, this jet is most likely substantially inclined $\\sim(15-25)^{\\circ}$ with respect to the observer, so that the broadband emission observed from M87 is not necessarily (everywhere) swamped by relativistically-beamed jet radiation. M87 thus appears to be a special active galaxy where we may be able to discover central features related to other high-energy processes. The rapid (days-scale) VHE flux variability detected during high source states by current $\\gamma$-ray instruments is the fastest variability seen so far at any wavelength, and suggests that the variable, hard-spectrum $\\gamma$-ray emission may originate close to its supermassive black hole. A variety of physical scenarios have been proposed in the literature to accommodate such findings. Some of the pros and cons/limitations of these candidate models are mentioned in this review, and the relevance of further VHE observations is stressed. VHE flaring characteristics (energy spectra and extension, variability time scales) like those reported bear the rare potential to give us important insights into the near black hole environment of active galaxies. Future, more sensitive VHE observations (CTA) will allow to study spectral variability and further constrain the time scales of VHE flux variations - both important inputs for theoretical modeling. Simultaneous VHE/high-resolution radio observations, on the other hand, will have the potential to pin down the location of the VHE emission on spatial scales $\\lppr 10^2 r_s$. M87 is the best-suited AGN for such kind of studies, since it is (beyond Sgr~A*) the most promising object for which direct imprints of the black hole in future high-resolution radio observations can be expected.\\cite{dexter12}"
    },
    "1208/1208.4130_arXiv.txt": {
        "abstract": "In the ultraviolet (UV), Type Ia supernovae (SNe Ia) show a much larger diversity in their properties than in the optical.   Using a stationary Monte-Carlo radiative transfer code, a grid of spectra at maximum light was created varying bolometric luminosity and the amount of metals in the outer layers of the SN ejecta.  This model grid is then compared to a sample of high-redshift SNe Ia in order to test whether the observed diversities can be explained by luminosity and metallicity changes alone.  The dispersion in broadband UV flux and colours at approximately constant optical spectrum can be readily matched by the model grid. In particular, the UV1-b colour is found to be a good tracer of metal content of the outer ejecta, which may in turn reflect on the metallicity of the SN progenitor. The models are less successful in reproducing other observed trends, such as the wavelengths of key UV features, which are dominated by reverse fluorescence photons from the optical, or intermediate band photometric indices. This can be explained in terms of the greater sensitivity of these detailed observables to modest changes in the relative abundances. Specifically, no single element is responsible for the observed trends. Due to their complex origin, these trends do not appear to be good indicators of either luminosity or metallicity. ",
        "introduction": "Type Ia supernovae (SNe Ia) are some of the most important tools for current cosmological studies.  Following the discovery that their peak magnitudes could be standardised \\citep{P93}, their use enabled the discovery that the universe is accelerating \\citep{R98,P99}.  More recent supernova observations, combined with other constraints from the cosmic microwave background and baryon acoustic oscillations have established that we live in a flat universe with a matter content of $\\Omega_M \\approx 0.27$ \\citep{Astier06,WoodVasey:2007ky,Kessler:2009et,Sullivan:2011bu} and the remaining 73\\%\\ made up of dark energy, the nature of which is currently unknown.  SNe Ia studies have measured that $w$, the dark energy equation of state parameter is consistent with $w = -1$ to 6.5\\% with some studies at high-z even beginning to place constraints on whether $w = w(z)$ \\citep{Riess07}.  In the future, we will have to observe at higher redshifts in order to find supernovae from younger times in the universe to improve on these dynamical measurements. As such our surveys will either have to switch to the IR or probe the rest-frame UV, a region of the spectrum that has been less extensively explored in the local population.  The UV spectra of SNe Ia have long been thought to probe the region where metallicity effects would be important \\citep{Hoflich:1998eg,Lentz:2000ff} due to the vast number of metal line transitions in this region.  Many of the photons in this region are absorbed, mostly by iron group elements, in what is referred to as metal line-blanketing; however, the effect that metallicity has on the level of the continuum flux in this region is debated \\citep{Sauer:2008gm}.  One example of the effect that progenitor metallicity may have is that in a higher metallicity progenitor, the production of neutron-rich isotopes such as $^{54}$Fe and $^{58}$Ni is favoured compared to \\nickel\\ \\citep{Iwamoto:1999}. This will be reflected not only in the spectra, but also to some degree in the broad-band light curves. A SN Ia in the local universe, where metallicity is high, will on average have a lower luminosity than a high-redshift SN Ia due to the different \\nickel\\ content and hence a different light curve stretch, as shown in \\citet{Timmes:2003bp,Howell:2007fl,Howell:2009fm}.  The abundance ratio of stable iron-group elements to radioactive \\nickel\\ has been proposed as an additional parameter for the standardisation of SN Ia light curves \\citep{Mazzali:2006cy}.  Until recently, our understanding of the role of progenitor metallicity in SNe\\,Ia data was limited by the paucity of observed UV spectra. Recognising this, \\citet[][hereafter, E08]{Ellis08} secured high quality Keck  spectra for 36 intermediate redshift ($z\\simeq$0.5) SNe\\,Ia at maximum light drawn from the Supernova Legacy Survey (SNLS); these optical spectra appropriately probe the rest-frame UV. \\defcitealias{Ellis08}{E08} \\citetalias{Ellis08} noted a significant diversity in their UV spectra which could not be attributed to dust. Importantly, they found the variations in their UV data, as characterised by colours derived directly from their rest-frame spectra, did  not correlate with the light-curve stretch. They also showed that the wavelengths of specific UV features showed phase-dependent shifts. The dispersion in their UV colours was claimed to be larger than could be accounted for metal-dependent models available at the time \\citep{Lentz:2000ff} thus opening the possibility of an additional explanation for the diversity.  Recently, Hubble Space Telescope (HST) and Swift observations have begun to explore the rest-frame UV of local SNe Ia.  \\citet{Foley:2008fx} used archival HST and International Ultraviolet Explorer (IUE) data to show that a particular ratio of UV flux correlates strongly with absolute V-band magnitude for 6 objects with spectra near maximum light: brighter supernovae have lower values of the ratio.  This is a different result from that claimed in \\citetalias{Ellis08} which saw no correlation with supernova brightness.  In a more recent study using 21 intermediate redshift SNe Ia ($z\\simeq$0.25), \\citet{Foley:2012ej} repeat the UV flux ratio analysis and find a different relation between absolute V-band magnitude and the ratio value.  In this case, brighter supernovae still show lower values of the ratio value, but the slope of the relation is very different. \\citet{Foley:2008fx} explored the use of the UV ratio as a luminosity indicator for light curve standardisation with some degree of success.  The potential link between UV properties and intrinsic luminosity would have important implications for future cosmological studies.  Swift data have been used by \\citet{Brown:2010cd} and \\citet{Milne:2010em} to obtain an overview of the spectral behaviour in the UV. In the near-UV filters (2600--3300\\AA\\ and 3000--4000\\AA), the normal sub-class of SNe Ia shows a high degree of homogeneity, while the subluminous and the peculiar SN\\,2002cx-like groups show large differences. Absolute magnitudes at maximum brightness are correlated with the optical decay rate and show a scatter similar in size to that obtained with optical data.  However, in the mid-UV (2000--2400\\AA) the scatter is much larger ($\\sim 1$\\,mag), indicating possible metallicity-driven effects in this part of the spectrum.  Recent HST observations of 12 Hubble-flow SNe Ia at maximum light by \\citet{Cooke:2011jm} show that the dispersion from a mean spectrum increases as wavelength decreases, and is largest in the UV region of the spectrum.  They attribute this to the larger number of metal absorption lines in the UV compared to the optical.   The same increase in dispersion is seen at higher redshift to the same degree \\citepalias{Ellis08} so they conclude that this must be an intrinsic feature of the supernova and not due to evolutionary effects.  A larger study of UV spectra at maximum from HST is underway (Maguire et al.~2012, in press).  A large degree of diversity is also seen in the UV photometry and spectra of four supernovae discussed in \\citet{Wang:2012cn}.  The paper concludes that more detailed modelling of supernovae in the UV is required.  Optical studies have shown that supernova properties depend on the properties of the host galaxy.  \\citet{Hamuy:2000il} first showed that the higher mass galaxies preferentially host dimmer supernovae compared to brighter supernovae which were associated with younger stellar populations in late-type galaxies.  \\citet{Sullivan:2010ez} examined SNe Ia subdividing their sample by host properties.  They found that in more massive galaxies, or in those with a lower specific star-formation rate, the SNe Ia were on average $\\simeq 0.08$\\,mag brighter than that of SNe Ia in other galaxies after correction for light-curve stretch and colour. \\citet{Sullivan:2010ez} suggested that the difference they observe may be due to the metallicity of the host galaxy as more massive galaxies tend to be more metal-rich; however this appears at odds with the results of \\citet {Timmes:2003bp} and \\citet{Mazzali:2006cy}.   This study is thus motivated by the need to reconcile the conflicting deductions regarding the observed diversity in the UV spectra derived from earlier work  We exploit a wide range of models parameterised by both bolometric luminosity and metallicity to see if we can explain the observations with these two variables alone.  The model dataset is presented in Section \\ref{sec:model-dataset} and the optical data sample which we use for comparison is described in Section \\ref{sec:obs-data}.  In Section \\ref{sec:analysis} we compare measurements of various UV properties of the model and data samples.  Our results are then discussed in Section \\ref{sec-discussion} and a summary of our conclusions presented in Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have used maximum-light spectra for a sample of 9 SNe Ia  obtained by \\citetalias{Ellis08} and compared them to a series of 1D models produced by a MC radiative transfer code in order to test whether the variations observed the UV region of the spectrum can be explained by metallicity and luminosity changes alone.  Our model spectra were initially based on one-well observed example SN\\,2005cf.  We summarise our conclusions in the following points  \\begin{itemize} \\item Our models replicate well the range of broad-band UV colours seen in the observed spectra, showing that the dispersion increases as one progresses to shorted wavelengths.  The (UV1-b) colour appears to depend strongly on metallicity, and it might have the potential for use in light-curve standardisation because it is less sensitive to very small changes in absorption features than narrow-band indicators(Sections \\ref{sec:disc_colours} and \\ref{sec:standarise}); however this requires further observational data for testing at lower luminosities. \\item We observe that $\\lambda_1$ and $\\lambda_2$ can be interpreted as peaks caused by the reverse fluorescence of photons into the UV region of the spectrum, mainly from intermediate mass elements and chromium (Section \\ref{sec:disc_l1l2}). Both $\\lambda_1$ and $\\lambda_2$ move towards the blue with increasing metal content in the upper layers of the ejecta, but the effect is highly non-linear and may not be the best metallicity indicator. \\item We see that $R_{UV}$ in our models has a low dependence on luminosity and the effect of the metal content is complex (Section \\ref{sec:disc_ruv}). High-z data do not confirm the relation of \\cite{Foley:2008fx}, and so we do not support using this index for light curve standardization. \\item We have used high-z Type Ia supernovae spectra from \\citepalias{Ellis08} and measured observed values of $R_{UV}$ for 9 objects around maximum light. We see that the measured values agree with those in the models, although they are affected by large errors from noise in the spectra (Figure \\ref{fig:r_uv}). \\item We have performed an extensive search using the model spectra in order to investigate the possibility that these results are driven by one or two elemental transitions, but the results have shown that the UV spectrum is far too complicated for this.  In light of this we suggest that the use of $R_{UV}$ is not appropriate for light-curve standardisation as the trends with metallicity and absolute magnitude are not linear (Section \\ref{sec:elements}). \\end{itemize}"
    },
    "1208/1208.1523_arXiv.txt": {
        "abstract": "A two-body system hypothetically affected by an additional radial acceleration $H v_r$, where $v_r$ is the radial velocity of the binary's proper orbital motion, would experience long-term temporal changes of both its semimajor axis $a$ and the eccentricity $e$ qualitatively different from any other standard competing effect for them. Contrary to what one might reasonably expect, the analytical expressions of such rates do not vanish in the limit $M\\rightarrow 0,$ where $M$ is the mass of the primary, being independent of it. This is a general requirement that any potentially viable physical mechanism able to provide such a putative acceleration should meet. Nonetheless, if $H$ had the same value $H_0$ of the Hubble parameter at present epoch, such rates of change would have magnitude close to the present-day level of accuracy in determining planetary orbital motions in our Solar System. However, general relativity, applied to a localized gravitationally bound binary system immersed in an expanding Friedmann-Lema\\^{\\i}tre-Robertson-Walker, does not predict the existence of such a putative radial acceleration at Newtonian level. Instead, it was recently shown in literature that an acceleration of order $H$ and directed along the velocity $\\bds v$ of the test particle occurs at post-Newtonian level. We worked out its orbital effects finding well-behaved secular rates of change for both $a$ and $e$ proportional to the Schwarzschild radius $r_s$ of the primary. Their magnitude is quite small: the rate of change of $a$ amounts to just 20 $\\mu$m per century in our Solar System. Finally, we discussed certain basic criteria of viability that modified models of gravity should generally  meet when their observable effects are calculated. ",
        "introduction": "In this paper, we first deal with a  certain hypothetical anomalous radial acceleration  proportional to the radial velocity of the orbital motion of a two-body system through a coefficient $H$ having dimensions of ${\\rm T}^{-1}$.  Some intuitive, Newtonian-like guesses about the possibility that such a putative acceleration may exist as a local manifestation of the cosmic expansion in the case of non-circular motions are offered in Section \\ref{conget}. They are motivated by the known fact that, within certain limits, several well established key features of a homogeneous and isotropic expanding universe, and also of its influence on local gravitationally bound systems, can be practically inferred within a classical framework \\citep{1965AnPhy..35..437H,1996MNRAS.282..206T,2010RvMP...82..169C,2012arXiv1203.5596B,2012arXiv1207.0060F}.  Such a putative extra-acceleration may, in principle,  have interesting phenomenological consequences  since it would induce peculiar orbital signatures which would not be mimicked by any other known competing dynamical effect. Moreover, by assuming for $H$ a value equal to that of the Hubble parameter at present epoch, the magnitude of these exotic effects for the planets of our Solar System would be close to the current level of accuracy in determining their orbits. Tensions might even occur between data and predictions in the case of Mercury and Mars. These topics are treated in Section \\ref{osser}  In Section \\ref{vaffa} we try-unsuccessfully-to find a theoretical justification for the guesses of Section \\ref{conget} rooted in a full general relativistic treatment of the orbital dynamics of a local system embedded in an expanding homogeneous and isotropic Friedmann-Lema\\^{\\i}tre-Robertson-Walker (FLRW) spacetime metric.  Nonetheless, our results remain valid from a phenomenological point of view because of their actual independence of any specific theoretical scheme, and can be viewed as observational constraints on such a putative exotic force, whatever the physical mechanism yielding it (if any) may be. Moreover, we feel that the numerical values coming out from our analysis  are interesting if compared with the observations, and may pursue further investigations to find a possible physical origin, cosmological or not, for it. Cosmological effects linear in $H$ were recently derived by Kopeikin \\citep{2012arXiv1207.3873K} for the propagation of electromagnetic waves between atomic clocks in geodesic motion in a FLRW background.  In Section \\ref{conside}  we build on the certain aspects discussed in previous Sections and provide some very general viability criteria that must be met by modified models of gravity \\citep{2012PhR...513....1C} in order not to give rise to unphysical observable effects. In particular, it is important to check the behaviour of their detectable predictions in the limits $G\\rightarrow 0,M\\rightarrow 0,$ where $G$ is the Newtonian constant of gravitation and $M$ is the mass of the central body acting as localized source of the gravitational field.  Section \\ref{conclusioni} provides an overview of the results obtained. ",
        "conclusions": "\\lb{conclusioni} We investigated the possibility that the cosmological expansion may impact the orbital motion of a localized two-body system at first order in the Hubble parameter $H$.  Reasoning classically, intuitive guesses  lead us to postulate, at Newtonian level, the existence of an additional radial acceleration $A^{(H\\rm N)}_r$ of order $H$ proportional to the radial velocity $v_r^{(\\rm orb)}$ of the proper motion of the test particle with respect to the primary. By considering $A^{(H\\rm N)}_r$ as a small correction to the Newtonian monopole $A^{(\\rm N)}$, we perturbatively worked out its long-term effects on the  Keplerian orbital elements of the test particle and on its distance from the primary, and, in the case of a binary in the plane of sky, on the projection of its orbit onto the line-of-sight and on its radial velocity. It turned out that both the semimajor axis $a$ and the eccentricity $e$ of the test particle would secularly increase. The analytical expressions of their rates of change are independent of the mass $M$ of the primary, so that they do not vanish in the limit $M\\rightarrow 0$, contrary to that one might reasonably expect. Such a feature constitutes a general requisite to be satisfied in the search of potentially viable alternative, i.e. non-cosmological, physical mechanisms able to provide $A_r^{(H\\rm N)}$.  Nonetheless,  the consequences of $A^{(H\\rm N)}_r$ would be interesting from a phenomenological point of view since their magnitude is close to the current level of accuracy in determining the planetary orbits in our Solar System.  Then, we looked for theoretical justifications of the existence of $A^{(H\\rm N)}_r$ in three different frameworks within general relativity. None of them provided it at Newtonian level, contrary to the known Hooke-like term of order $H^2$. Instead, as recently pointed out in literature, a velocity-dependent acceleration $A^{(H{\\rm p N})}$ of order $H$ and directed along the velocity $\\bds v$ of the test particle  exists at post-Newtonian level. We perturbatively worked out its orbital effects by finding secular rates of change of $a$ and $e$ proportional to the Schwarschild radius $r_s$ of the primary. They are well-behaved since they correctly vanish in the limit $M\\rightarrow 0$. For a planet of our Sun such effects are negligibly small: suffice it to say that the semimajor axis increases at a rate of just 20 $\\mu$m per century. We discussed the limits of validity of the orbital precession due to the $H^2$ term. We noticed that the formal singularity occurring in it for $M\\rightarrow 0$ actually has no physical meaning since such a limit is beyond the regime of validity of the perturbative calculation yielding the precession itself.  Finally, we extended some points emerged in dealing with the previous cosmological issues by discussing certain basic criteria of viability that \\textcolor{black}{certain classes of} modified models of gravity should generally  meet in view of their predicted observable effects like, e.g., orbital precessions. Given that such modified gravities must reduce to small perturbations of standard Newtonian gravity in appropriate circumstances, their precessions, which slowly alter an otherwise fixed Keplerian ellipse, must necessarily vanish in the limit of no gravity, i.e. for $G\\rightarrow 0$\\textcolor{black}{, at least as far as modified gravities whose perturbations to general relativity scale with $G$ are concerned; it is not necessarily valid in all cases}. The same should occur also for $M\\rightarrow 0$, unless such a limit violates the validity of the pertubative regime in which the precessions are calculated."
    },
    "1208/1208.0844_arXiv.txt": {
        "abstract": "The Keck Array (\\spud) is a set of microwave polarimeters that observes from the South Pole at degree angular scales in search of a signature of Inflation imprinted as B-mode polarization in the Cosmic Microwave Background (CMB).  The first three Keck Array receivers were deployed during the 2010-2011 Austral summer, followed by two new receivers in the 2011-2012 summer season, completing the full five-receiver array.  All five receivers are currently observing at 150~GHz. The Keck Array employs the field-proven \\bicep/\\bicepp strategy of using small, cold, on-axis refractive optics, providing excellent control of systematics while maintaining a large field of view.  This design allows for full characterization of far-field optical performance using microwave sources on the ground.  We describe our efforts to characterize the main beam shape and beam shape mismatch between co-located orthogonally-polarized detector pairs, and discuss the implications of measured differential beam parameters on temperature to polarization leakage in CMB analysis. ",
        "introduction": "\\label{sec:intro} Cosmological Inflation is a theory that describes the entire observable Universe as a microscopic volume that underwent violent, exponential expansion during the first fraction of a second.  Inflation is supported by the flatness and extreme uniformity of the Universe observed through measurements of the Cosmic Microwave Background (CMB)~\\cite{boomerang,wmapFlat}. Measurements of the polarization of the CMB could prove to be an impressive tool for probing the epoch of Inflation. A generic prediction of Inflation is the production of a Cosmic Gravitational-Wave Background, which in turn would imprint a faint but unique signature in the polarization of the CMB that has a curl component~\\cite{kamionkowski, seljak}.  This curl component of the polarization field is commonly called B-mode polarization, while the curl-free component, dominated by production due to density fluctuations at the time of last scattering, is called E-mode polarization.  % The strength of the B-mode polarization signature depends on the energy scale of Inflation, and would be detectable % if Inflation occurred near the energy scale at which the fundamental forces unify ($\\sim 10^{16}$~GeV).  The Keck Array, also called \\spud, is a set of five degree-scale microwave polarimeters that is currently observing the CMB from the Martin A. Pomerantz Observatory at the South Pole in search of a B-mode polarization signature from Inflation~\\cite{sheehy}. Each of the five receivers has 512~Transition Edge Sensor (TES) detectors, 16 of which are dark, leaving 496~detectors that are coupled to planar arrays of slot antennas, for a total of 2480~optical detectors in the entire instrument.   Receivers utilize a compact, on-axis refracting telescope design.  The first three Keck Array receivers were deployed during the 2010-2011 Austral summer and two new receivers followed in the 2011-2012 deployment season.  The first season of observation with the full five-receiver array is currently underway, with all receivers observing at 150~GHz.  The modular design of the Keck Array allows for future upgrades to include replacement of individual receivers to provide additional frequency coverage at 100~GHz and 220~GHz. The Keck Array leverages field-proven techniques employed for the \\bicep and \\bicepp telescopes, but with a vastly increased number of detectors, leading to increased sensitivity to the tiny Inflationary B-mode signal. The current upper limit on the B-mode amplitude in the CMB is set by the Keck Array's predecessor experiment, \\bicep~\\cite{chiang, takahashi}, and corresponds to $r < 0.72$ at 95\\% confidence level, where $r$ is the tensor-to-scalar ratio. The Keck Array aims to reach a sensitivity corresponding to $r=0.01$, where gravitational lensing of E-modes into B-modes should begin to be comparable in strength to the Inflationary signal.  As sensitivity dramatically improves with each generation of experiments, control of systematics becomes increasingly important.  Precise characterization of the optical performance of Keck Array receivers is critical to reach these ambitious sensitivity goals.  In the simplest mapmaking schemes, differential beam effects between co-located orthogonally-polarized pairs of detectors can lead to leakage of the CMB temperature signal into the much smaller B-mode signal, potentially limiting the ability of the Keck Array to reach its design sensitivity if these systematics are not well understood.  Characterizing the beam pattern of each of the 2480~Keck Array detectors in the far field presents a challenge in both data acquisition and reduction. We describe here our effort to characterize the Keck Array optical performance through an extensive ground-based precision beam mapping campaign at the South Pole, and discuss our understanding of the cause of measured beam non-idealities and our strategy for mitigating the effect of measured differential beam components in CMB analysis.  Four companion papers are also presented at this conference, focusing on the status of \\bicepp and the Keck Array (Ogburn {\\it et al.}~\\cite{ogburnSPIE}), the sensitivity of the Keck Array (Kernasovskiy {\\it et al.}~\\cite{kernasovskiySPIE}), the performance of the dual-polarization planar antenna array (O'Brient {\\it et al.}~\\cite{obrientSPIE}), and the thermal stability of \\bicepp (Kaufman~{\\it et~al.}\\cite{kaufmanSPIE}). ",
        "conclusions": "\\label{sec:conclusions} Pushing deeper into the level of B-mode polarization to constrain $r$ requires the dramatically increased sensitivity that the Keck Array provides, but also requires tight control of systematics. Through a massive beam mapping campaign, we have measured beam properties of each of the 2480~detectors in the Keck Array. We extract differential beam parameters using the same linear basis that will be employed for analysis mitigation of these modeled effects.  The source of the dominant pointing mismatch is still under investigation.  We believe that there is a complex relationship between the size and orientation of the near-field mismatch and that of the far field.  Measurements with new reduced near-field mismatch focal planes in the far field are upcoming.  We expect that the next generation of Keck Array focal planes will benefit from dramatically reduced far-field mismatch.  For the current receivers, we have shown that noise-dominated upper limits placed on the unmodeled residual component of the difference beam pattern already reach the $10^{-5}$ level at $\\ell=100$ when including one pair of detectors only.  We are optimistic that improved constraints on this residual, together with full simulations which account for averaging-down effects from observing the sky with many detectors at multiple drum angles, will demonstrate beam systematic control more than sufficient to reach $r=0.01$."
    },
    "1208/1208.0596_arXiv.txt": {
        "abstract": "{Characterization of the properties of young brown dwarfs are important to constraining the formation of objects at the extreme low-mass end of the initial mass function.  While young brown dwarfs share many properties with solar-mass T Tauri stars, differences may be used as tests of how the physics of accretion/outflow and disk chemistry/dissipation depend on the mass of the central object. This article summarizes the presentations and discussions during the splinter session on {\\em Disks, accretion and outflows of brown dwarfs} held at the CoolStars17 conference in Barcelona in June 2012. Recent results in the field of brown dwarf disks and outflows include the determination of brown dwarf disk masses and geometries based on Herschel far-IR photometry (70-160\\,$\\mu$m), accretion properties based on X-Shooter spectra, and new outflow detections in the very low-mass regime. } ",
        "introduction": "\\label{sect:intro}  The exploration of disks, accretion and outflows of young brown dwarfs (BD) plays an important role in developing our understanding of BD formation, planet formation, and the physics of circumstellar disks and outflows in general. It is related to fundamental open questions in stellar astronomy, such as: Do BDs form via the same path as stars? Can planets form around BDs? How do disks develop in a low-gravity, -temperature, and -radiation environment?.  It has been established in the last years that BDs and very low-mass stars (VLMS) at an age of a few Myrs resemble higher mass T~Tauri stars. They exhibit surface activity, such as cool spots (e.g., Joergens et al. 2003; Rodr\\'iguez-Ledesma et al. 2009) and coronal activity (e.g., Stelzer et al. 2006). BD/VLMS have disks with a similar fraction as stars detected at mid-IR (e.g., Comer\\'on et al. 2000; Natta \\& Testi 2001; Jayawardhana et al. 2003; Luhman et al. 2008) and far-IR/submm wavelengths (e.g., Klein et al. 2003; Scholz et al. 2006; Harvey et al. 2012), which are actively accreting (e.g., White \\& Basri 2003;  Mohanty et al. 2005; Herczeg \\& Hillenbrand 2008; Rigliaco et al. 2011), and often show signs of grain growth and crystallization (e.g., Apai et al. 2005; Pascucci et al. 2009). The finding that very young BDs rotate on average slower (e.g., Joergens \\& Guenther 2001; Joergens et al. 2003; Caballero et al. 2004) than their older counterparts (e.g., Bailer-Jones \\& Mundt 2001; Mohanty \\& Basri 2003) is indicative of a magnetic braking mechanism due to interaction with the disk.  Outflows from BD/VLMS have been observed in spectro-astrometry of forbidden emission lines (e.g., Whelan et al. 2005, 2009; Bacciotti et al. 2011; Joergens et al. 2012a), and in images of CO J=2-1 emission (Phan-Bao et al. 2008).  Many details of the properties of BD disks and outflows, however, are still unclear: What are the masses and sizes of BD/VLMS disks? How does grain evolution occur in these disks? Do planets form in circum-BD disks? What are the properties of BD/VLMS outflows and how does the outflow mechanism work at low mass-accretion rates? Are there any discontinuities of all these properties with mass indicating a different formation path of BDs compared to stars? Several new instruments that are ideally suited to study disks and outflows of faint objects (PACS/Herschel, X-Shooter/VLT, ALMA-early science) are producing currently their first results. At the splinter session on {\\em Disks, accretion and outflows of brown dwarfs} at the CoolStars17 conference (Barcelona, June 2012, www.mpia.de/homes/joergens/cs17splinter.html) leading scientists as well as young researchers presented and discussed recent results in the field of BD/VLMS disks and outflows, as outlined in the following. ",
        "conclusions": ""
    },
    "1208/1208.6009_arXiv.txt": {
        "abstract": "{ We present the results from nine years of optically monitoring the gravitationally lensed $z_{\\mathrm{QSO}}=0.658$ quasar RX~J1131$-$1231. The R-band light curves of the four individual images of the quasar were obtained using deconvolution photometry for a total of 707 epochs. Several sharp quasar variability features strongly constrain the time delays between the quasar images. Using three different numerical techniques, we measure these delays for all possible pairs of quasar images while always processing the four light curves simultaneously. For all three methods, the delays between the three close images A, B, and C are compatible with being 0, while we measure the delay of image D to be 91 days, with a fractional uncertainty of $1.5\\%$ ($1\\sigma$), including systematic errors. Our analysis of random and systematic errors accounts in a realistic way for the observed quasar variability, fluctuating microlensing magnification over a broad range of temporal scales, noise properties, and seasonal gaps. Finally, we find that our time-delay measurement methods yield compatible results when applied to subsets of the data. } ",
        "introduction": "Using the time delays between multiple images of gravitationally lensed sources to measure cosmological distances \\citep{Refsdal:1964vh} has several advantages: there is no need for any primary or secondary calibrator, and there are no effects from the intergalactic or interstellar medium. The method, originally proposed for gravitationally lensed supernovae, has so far exclusively been applied to quasars lensed in most cases by individual massive galaxies. Exceptions are SDSS J1004$+$4112 and J1029$+$2623, two quasars lensed by galaxy clusters, with long time delays \\citep{Fohlmeister:2008df, Fohlmeister:2013ji}. The quasar lens time-delay method is now recognized as a tool that complements other cosmological probes, in particular for constraining $H_{\\rm{0}}$ as well as the dark-energy equation-of-state parameter, $w$ \\citep[e.g.,][]{Suyu:2012tw, Linder:2011cs, Moustakas:2009wj}. In spite of its advantages, the method has long faced two severe limitations to its effectiveness in constraining cosmology.  First, time delays between the gravitationally lensed images of a quasar are hard to measure. Some claimed time-delay measurements turned out to be erroneous \\citep[see, for instance, the controversy around Q0957+561:][and references therein]{Vanderriest:1989uj, Press:1992jj, Schild:1997hf, Kundic:1997br}. Understandably, early light curves tended to be short and sparse, often too short to clearly demonstrate that \\emph{microlensing} variability was not interfering with their analysis. Microlensing is seen as an uncorrelated extrinsic variability in the quasar images, which results from the time-variable magnification created by stars in the lensing galaxy \\citep[e.g.,][]{Chang:1979dk, Schmidt:2010ce}. In the best cases, light curves spanned a few years \\citep[see, e.g.,][]{Wyrzykowski:2003wv, Hjorth:2002gz, Burud:2002gb, Burud:2002ge}. One consequence is that the numerical methods used to measure time delays from these light curves were exceedingly ``optimistic'' in their assumptions about extrinsic variability. More recent measurements with better data frequently yielded delays inconsistent with the error estimates of the earlier measurements.  Second, given the measured time delays and lens and quasar image astrometry, there is a famed degeneracy between the \\emph{time-delay distance}, which is a scale parameter inversely proportional to $H_{\\rm{0}}$, and the spatial distribution of mass responsible for the strong-lensing phenomenon. The delays constrain only a combination of this time-delay distance and the surface density of the lens near the images \\citep{Kochanek:2002in}. This can be overcome with independent constraints on the structure of the lens. \\citet{Suyu:2009ig, Suyu:2010fq} convincingly showed that it is possible to control the effects of model degeneracies for B1608+656 \\citep[][CLASS survey]{Myers:1995gz}, a quadruply imaged quasar with accurate radio time delays \\citep{Fassnacht:2002ig}. To do this, the authors combined (1) detailed HST images of the lensed quasar host galaxy, (2) a velocity-dispersion measurement of the lens galaxy, and (3) information about the contribution of intervening galaxies along the line of sight, from galaxy number counts calibrated with numerical simulations.  In parallel to the advances in lens modeling, the observational situation has impressively evolved as well. Two observational groups, the COSmological MOnitoring of GRAvItational Lenses (COSMOGRAIL) and \\citet{Kochanek:2006fp}, have been intensely monitoring $\\approx 20$ lenses for roughly ten years. In 2010, our two groups decided to merge their observational efforts, with the COSMOGRAIL group focusing on the analysis of time delays and the \\citet{Kochanek:2006fp} group focusing on the analysis of microlensing. While preliminary results have been published both before and after this merger \\citep{Kochanek:2006fp,Vuissoz:2007du,Vuissoz:2008bu,Morgan:2008fc,Morgan:2008jb}, exquisite data spanning almost a decade of continuous observation are now being released, for instance for the quadruply imaged quasar HE~0435-1223 \\citep{Courbin:2011bl, Blackburne:2011ue}.   In this paper, we present nine years of optical monitoring of the quadruply imaged quasar RX~J1131$-$1231 \\citep{Sluse:2003cw}, and measure its time delays with the techniques of \\citet{pycs}. RX~J1131$-$1231 is one of the most spectacular lenses of our sample. The redshift of the lensing galaxy is $z_{\\mathrm{lens}} = 0.295$, while the quasar is at $z_{\\mathrm{QSO}} = 0.658$. This low quasar redshift means (1) that the photometric variations are fast, numerous, and \\emph{strong} because it is a lower-luminosity quasar \\citep[see, e.g.,][]{MacLeod:2010ex}, and (2) that the host galaxy of the lensed quasar is seen as a full Einstein ring with many spatially resolved structures in HST images. Similarly, the lensing galaxy is sufficiently bright to allow a precise measurement of its velocity dispersion and possibly of its velocity dispersion profile. These characteristics facilitate both the time-delay measurement and the lens modeling. The latter, with state-of-the-art inferences of cosmological constraints based on our time-delay measurements of RX~J1131$-$1231, are presented in \\citet{suyu1131}.  Our observations of RX~J1131$-$1231 and their reduction are described in Sections~\\ref{obs} and \\ref{reduc}, while the light curves are presented in Section \\ref{lc}. In Section~\\ref{timedelay}, we apply three different curve-shifting techniques to the light curves and infer our best measurements of the delays along with realistic random and systematic error bars. Our results are summarized in Section~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The first part of this paper describes the COSMOGRAIL data reduction procedure, which will be used to reduce all data gathered by our monitoring campaign of gravitationally lensed quasars. In the second part we apply this pipeline to our COSMOGRAIL and SMARTS observations of the quad lens RX~J1131$-$1231, leading to an unprecedented set of nine-year-long light curves of high photometric quality. Several strong and fast intrinsic quasar variability patterns constrain the time delays between the multiple images. Microlensing-related extrinsic variability is clearly present, as pointed out and analyzed in previous studies \\citep[][]{Sluse:2006jv, Sluse:2007hw, Morgan:2010bw,  Dai:2010im, Chartas:2012dl}. However, this distorting signal does not prevent us from measuring accurate time delays, using the three independent algorithms of \\citet{pycs}.  The best time-delay estimates of RX~J1131$-$1231 are provided by the regression difference technique. It measures the 91-day delays between D and the other quasar images to a fractional $1\\sigma$ uncertainty of $1.5\\%$. This error estimate is obtained by applying the techniques to synthetic curves with known time delays, which contain extrinsic variability features similar to the observed ones. We demonstrate the consistency of our error estimates by independently measuring time delays -- including error bars -- from subsets of the observed light curves of RX~J1131$-$1231. This experiment also reveals that long multi-year monitoring is essential for reliably measuring time delays, despite progress on the methods.  The results from this paper are used to constrain the time-delay distance toward RX~J1131$-$1231 and deduce stringent implications for cosmology in \\citet{suyu1131}. With our time-delay measurement errors of only 1.5\\%, the accuracy of this cosmological inference is actually limited by the residual uncertainty of (1) the gravitational potential of the lens galaxy and (2) the large-scale structures along the line of sight to the quasar \\citep{BarKana:1996gp}."
    },
    "1208/1208.1135_arXiv.txt": {
        "abstract": "Based on Bremer et al. (2011) and Eckart et al. (2012) we report on simultaneous observations and modeling of the millimeter, near-infrared, and X-ray flare emission of the source Sagittarius A* (SgrA*) associated with the super-massive (4$\\times$10$^6$\\solm) black hole at the Galactic Center. We study physical processes giving rise to the variable emission of SgrA* from the radio to the X-ray domain. To explain the statistics of the observed variability of the (sub-)mm spectrum of SgrA*, we use a sample of simultaneous NIR/X-ray flare peaks and model the flares using a synchrotron and SSC mechanism. The observations reveal flaring activity in all wavelength bands that can be modeled as the signal from adiabatically expanding synchrotron self-Compton (SSC) components. The model parameters suggest that either the adiabatically expanding source components have a bulk motion larger than v$_{exp}$ or the expanding material contributes to a corona or disk, confined to the immediate surroundings of SgrA*. For the bulk of the synchrotron and SSC models, we find synchrotron turnover frequencies in the range 300-400~GHz. For the pure synchrotron models this results in densities of relativistic particles of the order of 10$^{6.5}$cm$^{-3}$ and for the SSC models, the median densities are about one order of magnitude higher. However, to obtain a realistic description of the frequency-dependent variability amplitude of SgrA*, models with higher turnover frequencies and even higher densities are required. We discuss the results in the framework of possible deviations from equilibrium between particle and magnetic field energy. We also summarize alternative models to explain the broad-band variability of SgrA*. ",
        "introduction": "\\label{introduction}  Sagittarius~A* (SgrA*) is the closest super-massive black hole (SMBH) and is the prime candidate to study the variability and spectral properties of accreting SMBHs. Since it is at the center of the Milky Way at a distance of about 8~kpc, Doeleman et al. (2008, 2009) and  Fish et al. (2011) have recently succeeded in obtaining structural information on event-horizon scales through very long baseline interferometry (VLBI) at 1.3~mm wavelength. SgrA* is strongly variable in the radio and millimeter wavelength regime (Zhao et al. 2003; Mauerhan et al. 2005; Eckart et al. 2008ac; Marrone et al. 2008; Li et al. 2009; Yusef-Zadeh et al. 2008, 2009). It shows an inverted spectrum from the radio to the (sub-)millimeter domain (Falcke et al. 2000) and displays order-of-magnitude flares in the infrared and X-ray domain (Baganoff et al. 2001, 2003; Genzel et al. 2003; Eckart et al. 2004, 2006a, 2010, Ghez et al. 2003, 2004). The sub-mm spectrum of SgrA* itself is rather unexplored owing to the difficulty of separating it from contributions of the surrounding ''mini-spiral''  and the circum-nuclear disk (CND).  While the so-called sub-millimeter bump is attributed to SgrA* and is thought to be due to relativistic, thermal electrons of a hot, thick, advection-dominated accretion flow (Dexter et al. 2009, 2010; Narayan et al. 1995; Yuan et al. 2003), it is, however, currently unclear whether the radio spectrum is due to contributions from  an accretion flow (Yuan et al. 2003), due to a jet (Falcke \\& Markoff 2000), or a combination of both. The strong variability observed at radio to X-ray emission is most likely due to synchrotron and/or synchrotron self-Compton (SSC) radiation. Part of the radiation may be due to a single hot-spot or a multi-spot model in the mid-plane of the accretion flow or an increased accretion rate (see models by Broderick \\& Loeb 2006 and Eckart et al. 2006b, 2011, Meyer et al. 2007, Yuan et al. 2008, see also Pech\\'a\\v{c}ek et al. 2008, Shcherbakov \\& Baganoff 2010, Dexter et al. 2010, Zamaninasab et al. 2011). Spiral arm models (e.g. Karas et al. 2007) or jet/jet-base models (Falcke \\& Markoff 2000, Markoff 2005, Markoff et al. 2007) are also possible. Another result from multi-wavelength observations of  SgrA* is that the process of adiabatic expansion of source components may be relevant (Eckart et al. 2006a, 2008abc, 2009, 2010, 2012, Yusef-Zadeh et al. 2006a, 2008, 2009). Here we assume that the sub-mm, NIR and X-ray  flux excursions are physically related (However, see a detailed discussion on this in section 5.3 in Eckart et al. 2012). This expansion can explain the observed time lags between the infrared/X-ray and millimeter emission peaks. The fact that one observes different time delays between NIR and sub-mm flares can be understood as a consequence of a spread in source sizes and synchrotron turnover frequencies. The idea of adiabatic expansion is supported by the fact that there is no conclusive observational evidence for sub-mm flares preceeding NIR events and by the fact that it has been detected at radio cm-wavelengths (Yusef-Zadeh et al. 2006ab). Here we summarize results obtained for simultaneous flare emissions in the NIR and X-ray with implications to the observed radio variability. In this summary we include the analysis of the NIR spectral index under the assumption that especially for weak flares synchrotron losses are important. A detailed description of the Synchrotron/SSC analysis is given in Eckart et al. (2012). Throughout the article we use $S_{\\nu} \\propto \\nu^{-\\alpha}$. ",
        "conclusions": "\\label{summary} Eckart et al. (2012) have shown that the bulk of synchrotron and SSC models applicable to SgrA* have preferred synchrotron turnover frequencies in the range of a few hundred GHz. For the pure synchrotron models, the densities of relativistic particles are of the order of 10$^{6.5}$cm$^{-3}$. For models involving SSC contributions the median densities are an order of magnitude higher. These values are quite comparable to those quoted for the accretion flow toward SgrA* (e.g.  Yuan et al. 2003, 2004, Yusef-Zadeh et al. 2006ab).  For a realistic description of the observed frequency-dependent variability amplitudes of SgrA* higher densities and turnover frequencies are also required (details in Eckart et al. 2012). Higher densities are neither unreasonable nor out of reach. The accretion rate and hence the central density may be much higher than the limits derived from Faraday rotation (Bower et al. 2003, Bower 2003, Marrone 2006a, Marrone et al. 2006abc, and discussion by Marrone et al. 2007). In particular, the magnetic field equipartition fraction as well as the bias field strength in the case of magnetic field reversals - that can be expected for a turbulent flow - are unknown. Variations in the magnetic field structure as well as the field strength with respect to equipartition (Marrone et al. 2007, Igumenshchev et al. 2003) are likely to result in higher densities in the immediate vicinity of the central SMBH. While SgrA* may be close to equipartition, it is quite likely that at least during flux density events associated with very strong X-ray flares and the related NIR and (sub)mm-flux excursions, SgrA* may go through phases in which it is off equipartition.  Currently, it is difficult to decide which of all the models mentioned here describe the dominant contributors to the flare activity of SgrA*, or which combinations of effects are most likely or if one or some of the mechanisms can be excluded entirely. A key appears to be the restrictions imposed onto the models by the frequency dependant flux density variability or polarization properties. Therefore, the simultaneous flux density and polarization monitoring observing campaigns are essential. Any model (or combinations of those) must be able to explain these observations. In addition, it is essential, that high angular resolution observations in the NIR (interferometry with long baselines) and radio domain (mm- and sub-mm-VLBI) are being performed. These observations will help to identify the orbital motion around the SMBH or motion along a jet/outflow. \\\\ \\\\ {\\bf Acknowledgements: } N. Sabha is member of the Bonn Cologne Graduate School (BCGS) for Physics and Astronomy supported by the Deutsche Forschungsgemeinschaft. M. Valencia-S. is member of the International Max-Planck Research School (IMPRS) for Astronomy and Astrophysics at the Universities of Bonn and Cologne supported by the Max Planck Society. F. Baganoff was supported by NASA through Chandra Award Number GO9-0101X and SAO Award Number 2834-MIT-SAO-4018. Part of this work was supported by the COST Action MP0905: Black Holes in a violent Universe and PECS project No. 98040. We are grateful to all members of the NAOS/CONICA, ESO PARANAL, and APEX team. M. Garc\\'{\\i}a-Mar\\'{\\i}n is supported by the German federal department for education and research (BMBF) under the project number 50OS1101. R. Sch\\\"odel acknowledges support by the Ram\\'on y Cajal program, by grants AYA2010-17631 and AYA2009-13036 of the Spanish Ministry of Science and Innovation, and by grant P08-TIC-4075 of the Junta de Andaluc\\'ia. Ongoing CARMA development and operations are supported by the National Science Foundation under a cooperative agreement, and by the CARMA partner universities.  \\vspace*{0.5cm}"
    },
    "1208/1208.1904_arXiv.txt": {
        "abstract": "\\noindent After a careful analysis of the instrumental effects on the Poisson noise to demonstrate the feasibility of detailed stochastic variability studies with the \\textit{Swift} X-Ray Telescope (XRT), we analyze the variability of the black hole X-ray binary SWIFT J1753.5-0127 in all XRT observations during 2005-2010. We present the evolution of the power spectral components along the outburst in two energy bands: soft (0.5--2 keV) and hard (2--10 keV), and in the hard band find results consistent with those from the \\textit{Rossi X-ray Timing Explorer} (RXTE). The advantage of the XRT is that we can also explore the soft band not covered by RXTE. The source has previously been suggested to host an accretion disk extending down to close to the black hole in the low hard state, and to show low frequency variability in the soft band intrinsic to this disk. Our results are consistent with this, with at low intensities stronger low-frequency variability in the soft than in the hard band. From our analysis we are able to present the first measurements of the soft band variability in the peak of the outburst. We find the soft band to be less variable than the hard band, especially at high frequencies, opposite to what is seen at low intensity. Both results can be explained within the framework of a simple two emission-region model where the hot flow is more variable in the peak of the outburst and the disk is more variable at low intensities.\\\\ ",
        "introduction": "\\noindent Transient stellar-mass black hole candidate X-ray binaries (BHBs) in outburst broadly exhibit two different states: low luminosity hard states in which the spectrum is dominated by hard power law emission (out to a few tens of keV) from a hot inner flow, inner corona, and/or the base of a jet, and high luminosity soft states in which the spectrum is dominated by black-body emission (peaking at up to a few keV) from the accretion disk, often accompanied by a power law tail.  There is no agreement about the exact structure and origin of the power law emitting region, nor on the disk geometry: while previously it was thought that in the hard state the disk is truncated at a large distance of at least several tens of gravitational radii, recently it has been debated if instead it extends all the way down to the innermost stable circular orbit (ISCO). See \\citet{belloni2005}, \\citet{klis2006} and \\citet{remillard2006} for detailed descriptions of the different states along an outburst and \\citet{done2007} for a detailed discussion of the different models. \\\\ \\\\ The  ubiquitous X-ray variability in BHBs could provide useful constraints on the structure and behavior of the emitting regions, beyond that provided by the X-ray spectrum.  \\citet{lyub1997} proposed a model which is becoming widely accepted (e.g. \\citealt{chu2001,uttley2005,done2007}), where the variability originates as mass-accretion fluctuations in the accretion flow, which propagate towards and through the X-ray emitting regions to produce the observed flux variability.  Since the variability in the hard states is strong (few tens of percent fractional rms) while in the soft states it is much weaker (a few percent),  \\citet{chu2001} suggested that the hot, power law emitting inner flow is responsible for this variability, while the cool disk is stable.   Furthermore, they showed that in a `standard' geometrically-thin cool disk, high frequency fluctuations can survive only if they arise at small radii, as at larger radii, viscous damping will prevent them from propagating inward. In the hard state, the power spectrum consists of broad band-limited noise accompanied by Quasi Periodic Oscillations (QPOs). The low frequency break in the noise power spectrum could be associated with the transition radius to the hot flow at the inner edge of the cool disk, the QPO could be the Lense-Thirring precession of the hot flow, and the high frequency break may be associated with the innermost radius of the hot flow (\\citeauthor{stella1998} 1998, \\citeauthor{chu2001} 2001).\\\\ \\\\ The \\textit{Rossi X-ray Timing Explorer} (RXTE) Proportional Counter Array (PCA) was sensitive in the $>$ 2 keV band dominated by the power law emission. Most variability studies until recently were done with RXTE, so the evolution of the variability below 2 keV remained unexplored. Recent work by \\cite{wilkinson} and \\cite{uttley2011} using \\textit{XMM-Newton} showed that in the hard state of the BHBs SWIFT J1753.5-0127, GX~339-4 and Cyg~X-1 the disk contributes significantly to variability in the soft band ($<$ 2 keV) at Fourier frequencies below 1~Hz. Moreover, the disk variations precede the correlated power law variations by a few tenths of a second. These results cast doubt on the notion of a stable disk not contributing to the source variability, implying that (at least in the hard state and at low frequencies) substantial flux variability is contributed by the cool disk. This highlights the importance of exploring the variability in the soft band.  Like \\textit{XMM-Newton}, the \\textit{Swift} X-ray Telescope (XRT) has a lower energy bound (0.3 keV) than the RXTE PCA, providing another opportunity to explore the variability at lower energies dominated by the disk emission. This directly addresses the question of the origin of the variability. \\\\  \\begin{figure} \\includegraphics[width=6.3cm,height=8.8cm,angle = -90]{lc-swift.ps} \\includegraphics[width=6.05cm,angle = -90]{hc.ps} \\caption{Light curve of the outburst during 2005-2010 (top) and corresponding hardness-intensity diagram (bottom) in the energy bands indicated. Each point represents one observation. } \\label{lchc} \\end{figure} \\noindent SWIFT J1753.5-0127 was discovered with the \\textit{Swift} Burst Alert Telescope on 2005 May 30 \\citep{palmer2005} and has been in outburst since then. The source remained in the hard state (\\citeauthor{miller2006} 2006 and \\citeauthor{cadolle} 2007, \\citeauthor{zhang2007} 2007, \\citeauthor{ramadevi} 2007, \\citeauthor{chiang2010} 2010, henceforth  C07, Z07, RS07 and Ch10, respectively) until June 2009 (MJD~55010) but then was in a softer state until at least July 2010 (MJD~55404) (\\citealt{soleri2012}).  During 2005-06, RXTE observations show white noise (frequency independent) at low (0.2--0.4~Hz) frequencies, a QPO ($\\sim$0.7 Hz) and red noise (power decreases with frequency) at high (1.5--2~Hz) frequencies, which is typical of the hard state (Z07, RS07). Spectral analyses with \\textit{XMM-Newton} and RXTE suggested the presence of a cool accretion disk extending down to near the ISCO that could be modelled with a disk black-body component in addition to the power law (\\citeauthor{miller2006} 2006, \\citeauthor{reis2010} 2010, Ch10). This is at variance with the picture in which the disk is truncated at larger radii in the hard state. There are also arguments against the presence of a soft disk-like component (see, e.g., \\citeauthor{done2007} 2007 and \\citeauthor{hiemstra2009} 2009), which, however, do not refute the evidence for intrinsic disk variability found by \\citet{wilkinson}.\\\\ \\\\ In this work, we use the \\textit{Swift} XRT to study the energy-dependent behavior of the variability  of the BHB SWIFT~J1753.5-0127. While some XRT power spectra of BHBs were  previously reported (e.g., \\citealt{kennea1659}, \\citealt{curran1752-swift}), the present work is the first study of the energy dependence extending below 2~keV of all power spectral components in a BHB along an outburst.  As a prerequisite to this work, we study the instrumental effects on the Poisson-noise spectrum (see Appendix) and demonstrate that detailed variability studies are feasible with \\textit{Swift} XRT. In Section \\ref{obdata}, we describe the observations and data analysis. In Section \\ref{lchd}, we discuss the evolution of the outburst using light curves and the hardness intensity diagram. In Section \\ref{timeanalysis}, we first investigate the energy dependence of the variability in a model-independent way, and then fit Lorentzians to the power spectra of all observations to study the evolution of the components and their energy dependence.  We also study the energy dependence of the power spectra inside the 0.5--2 keV band. Finally, in Section \\ref{correlations} we study the correlations between the variability parameters and source intensity. In Section \\ref{discussion}, we discuss the interpretation of our results, and their model implications. ",
        "conclusions": "\\label{discussion} \\noindent We report on the energy dependent variability of the first six years of the very long and still ongoing outburst of the BHB SWIFT J1753.5-0127. We observe the Fast Rise Exponential Decay type light curve also reported in earlier works (C07, RS07, Z07 and Ch10). The HID does not resemble in either \\textit{Swift} or RXTE (\\citealt{soleri2012}) data to the typical \"q\"-shaped track followed by many other BHBs. The source remained in hard states throughout the outburst (C07, RS07, Z07, Ch10), until a transition to a softer state occurred in 2009 (\\citealt{soleri2012}), leading to our softest observations in 2010. As we performed very detailed power spectral studies with the \\textit{Swift} XRT, we first examined the behavior of the Poisson level and find it to be affected by instrumental effects (see the Appendix for details). We carefully  compared our results with the earlier RXTE results using the 2--10 keV band, and found them to be consistent.  \\\\ \\\\ The \"type-C\" QPO (peaked low frequency components accompanied by strong broadband noise in the intermediate states, e.g. \\citealt{casella2005}), the  break frequency, $\\nu_b$ and the hump frequency, $\\nu_h$, reported by \\citealt{soleri2012} correspond to the QPO, the 0.1 Hz and 1 Hz components, respectively, in our analyses. We report  these components,  for the first time, in soft X-rays together with the harder band. The WK correlation, which is a positive correlation between  $\\nu_b$  and $\\nu_h$  \\citep{wk} was reported by \\citealt{soleri2012} between these components in SWIFT J1753.5-0127.  As there are very few simultaneous detections of these components in our data, we do not see a clear WK correlation in either band  and is hence not shown here. The fractional amplitudes of all three components vary over similar range (5-18 \\%) in both bands, except the 1 Hz component in the hard band which is the strongest (20-25 \\%) amongst all components in the peak of the outburst. In the decay, all the components get stronger in both energy bands, with no detections of the 1 Hz component. The analysis in soft sub-bands and hard band (see Figure \\ref{pds2005} and Table \\ref{tab:combob}) in the peak of the outburst shows that the 0.1 Hz component displays a flat rms spectrum i.e. the fractional amplitudes do not vary with energy. The QPO and the 1 Hz component cannot be constrained in the 0.5--1 keV band. These get stronger at higher energies, with the QPO getting stronger above 2 keV while the 1 Hz component gets stronger above 1.5 keV.    \\\\ \\\\ The main new finding of our work (Figure \\ref{rms}) is that in the peak of the outburst (for $>$120~c/s) the variability is weaker in the soft band  than in the hard band; the difference is most pronounced at high ($>$0.4~Hz) frequency but also seen at low ($<$0.4~Hz) frequency. Previous work with \\textit{XMM-Newton} \\citep{wilkinson} had shown that in the decay, in 2006, the opposite was true: at low frequency the soft band was more variable than the hard band, while at higher frequencies both bands had similar fractional variability amplitudes. Our 2005-2008 results in the decay are consistent with this.  \\citet{wilkinson} interpreted this extra soft-band variability in terms of a stronger contribution of the soft disk emission to the variability at low frequencies, which could be associated with accretion fluctuations intrinsic to the disk at these frequencies.  Combining our new results on the outburst-peak with these earlier findings on the decay, we can say that in both frequency ranges the fractional rms spectrum becomes softer as the source decays, while instead the overall source spectrum becomes harder.\\\\ \\\\ These findings can be understood in terms of a simple two-component representation (hot inner flow and cool disk) of the picture of \\citet{lyub1997} and \\citet{chu2001}.  The softer overall spectrum in the peak of the outburst is due to a relatively stronger disk contribution. That the fractional rms spectrum gets harder at all frequencies when the disk contribution increases can be explained if the observed disk variability has a lower fractional rms than the hot flow: the increased disk flux dilutes the variability at low energy.  This would imply that while in the decay the disk at low frequencies  is more variable (in fractional rms) than the hot flow, in the outburst-peak it is {\\it less} variable.  So, we may be seeing in this outburst-peak a transition to the less variable disk state that must also characterize the soft states. That in our outburst-peak data the difference between low and high energy variability is less at low frequency can be due to a residual contribution of intrinsic disk variability at low frequency similar to that proposed by \\citet{wilkinson} and \\citet{uttley2011}. \\\\ \\\\ It would be interesting to observe the variability behavior of this source when (if) it makes a transition back to the hard states, to see if it exhibits behavior similar to other BHBs during this transition, even if up to now it did not follow the usual \"q\" path in the HID. With our results, we demonstrate that the power spectral studies can be successfully performed with \\textit{Swift} XRT. The soft band provides very useful insight into the energy dependence of the variability which is essential to understand the contribution of the accretion disk. To understand better the role of the disk, variability studies of more BHBs with \\textit{Swift} XRT are required. We recommend caution while dealing with the  Poisson spectrum of the power spectra generated with \\textit{Swift} XRT."
    },
    "1208/1208.1303_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro} Active galactic nuclei (AGN) with relatively small black hole (BH) masses ($\\mbh \\lesssim 10^{6}~M_{\\odot}$) comprise a demographic that until recently went relatively unexplored.  It is difficult to find and observe such low-luminosity objects \\citep{2003ApJ...588L..13F, 2004ApJ...607...90B}.  While luminous quasars with massive central BHs have been studied intensively for decades \\cite[e.g.][]{1992ApJS...80..109B}, only with the advent of large surveys like the Sloan Digital Sky Survey \\citep[SDSS;][]{2000AJ....120.1579Y} has it become possible to search for large numbers of low-mass active galaxies \\citep{2004ApJ...610..722G, 2007ApJ...670...92G}.  The number density and radiative properties of these accreting, low-mass BHs, which we refer to as low-mass AGN, are vital to understanding the accretion history of the Universe.  By understanding their radiative properties, we can constrain the spectral energy distributions used in models of primordial seed BHs \\citep{2009MNRAS.400.1911V}, which may have prompted early galaxy formation \\citep[e.g.][and references therein]{2011ARA&A..49..373B} and were possibly a significant participant in the reionization of the Universe \\citep{2009ApJ...696L.146M}.  In addition, investigating these low-mass AGN is necessary to understand the overall demographics of BHs at low masses, including the low-mass end of the $M_{\\rm BH}-\\sigma^{\\ast}$ relation, which can further inform our understanding of galaxy formation and evolution \\citep{2002ApJ...574..740T, 2009ApJ...698..198G}.  We focus on the sample of 174 low-mass AGN selected by \\citet[][; hereafter GH07]{2007ApJ...670...92G} from the SDSS.  Note that these objects can be technically classified as narrow-line Seyfert 1 (NLS1) galaxies based on their broad H$\\beta$ widths \\citep{1985ApJ...297..166O}.  These low-mass AGN and their host galaxies are well-studied, with stellar velocity dispersion measurements presented in \\cite{2005ApJ...619L.151B} and \\cite{2011ApJ...739...28X}, and host galaxy analysis with the \\textit{Hubble Space Telescope} presented by \\cite{2008ApJ...688..159G} and \\cite{2011ApJ...742...68J}.  In addition, their spectral slopes from the UV to the X-ray have been used to assess whether the spectral energy distribution gets harder when the BH mass is lower \\citep[][Dong et al. in prep]{2007ApJ...656...84G, 2009ApJ...698.1515D}.  In this paper we use high-resolution spectroscopy from the Keck telescope to specifically study the emission-line properties of 27 of these objects. The Keck spectra are observed with a smaller aperture than the SDSS data, using a 0.75\" x 1\" aperture extraction versus 3\" fibers used in SDSS, with $\\sigma\\sim 22~$km~s$^{-1}$ instrumental resolution. This improves our ability to distinguish emission from the narrow-line region (NLR) of the AGN compared to \\ion{H}{2} regions in the host galaxy and makes it possible to decompose the broad and narrow components of the permitted lines more clearly than is possible in the SDSS data.  We make comparisons between AGN with low- and high-mass BHs, and investigate the metallicity of the NLR, the structure of the NLR, and the radiative properties of the low-mass AGN themselves.  In Sections \\S\\ref{sec:sample} and \\S\\ref{sec:method} we present our high resolution, high signal-to-noise (S/N) observations of a sample of 27 low-mass AGN and our fitting methodology and measurements.  In Section \\S\\ref{sec:metal}, we investigate the gas-phase metallicities of these AGN using various emission line diagnostics in pursuit of rare, low-metallicity AGN.  In Section \\S\\ref{sec:SED}, we discuss the inferred far-UV continuum slope and discuss implications for the spectral energy distributions of accreting low-mass BHs in AGN.  In Section \\S\\ref{sec:comps}, we investigate low-contrast, weak emission lines through use of composite spectra.  In particular, we focus on high-ionization Fe lines and their behavior relative to more prominent, lower-ionization NLR emission lines.  In Section \\S\\ref{sec:concl}, we present a summary of our conclusions. ",
        "conclusions": "\\label{sec:concl}  We present observations of a sample of 27 low-mass AGN ($10^{4}~M_{\\odot}<\\mbh<2~\\times~10^6~M_{\\odot}$), observed with the Echellette Spectrograph and Imager on the Keck Telescope.  Large samples of low-mass AGN have not existed until recently \\citep{2004ApJ...610..722G}, and so we compare their emission-line properties to those of well-studied higher-mass AGN.  We investigate the NLR metallicities of these objects using emission line ratios, particularly [\\ion{N}{2}]$/$H$\\alpha$ and [\\ion{O}{3}]$/$H$\\beta$.  Our sample includes objects with weaker [\\ion{N}{2}]$/$H$\\alpha$ than higher-mass AGN for a given [\\ion{O}{3}]$/$H$\\beta$, which implies that those objects have lower metallicities.  Thus we see some galaxies with similar metallicities to the rare AGN with sub-solar metallicity sought by \\cite{2006MNRAS.371.1559G}, yet we cannot determine their metallicities on an absolute scale \\citep[e.g.][]{2008ApJ...681.1183K}.  Additionally, we do see weak evidence for a correlation between galaxy stellar mass and gas-phase metallicity in these systems.  Most likely, the lack of low-metallicity AGN highlighted by \\cite{2006MNRAS.371.1559G} is attributable to the difficulties of finding AGN in low-mass galaxies \\citep[e.g.][]{2007ApJ...670...92G} rather than extra enrichment of the NLR by the AGN itself.  We examine the continuum properties of the accretion disk for our low-mass AGN.  By using the emission from the recombination lines [\\ion{He}{2}] and H$\\beta$, we infer the slope of the far-UV continuum from 200--900\\AA{}, near the peak of the big blue bump.  We find that the far-UV slopes in our low-mass AGN are steeper than those of low-redshift, higher-mass AGN.  While we know that our low-mass AGN are very radio-quiet \\citep{2006ApJ...636...56G} and have flat X-ray slopes \\citep{2009ApJ...698.1515D} compared to typical AGN, they have quite similar optical properties and unexpectedly, somewhat steeper UV slopes. These steeper inferred UV continua imply that changing continuum shape could explain the inverse Baldwin effect seen between H$\\beta$ FWHM and luminosity \\citep{2002MNRAS.337..275C, 2005ApJ...630..122G}.  Despite tentative evidence for different SEDs, the NLR structure is also similar between low-mass and typical AGN. Using composite spectra, we are able to measure the widths and velocities of high-ionization Fe lines. As found by \\cite{2009MNRAS.394L..16M}, the [\\ion{Fe}{7}] line exhibits a similar width and blue shift as the blue wing of [\\ion{O}{3}], pointing to a common physical origin for both transitions in a radially flowing component.  As seen in previous work, we find that the width of NLR lines and their blueshift correlates with the ionization potential of the line.  We see no evidence for a dramatic change in the NLR structure in this mass and luminosity regime.  Making composite spectra in subsets of luminosity, Eddington ratio, and presence or absence of a blue wing in [\\ion{O}{3}], we find first that the high-luminosity composite has much broader NLR emission lines than the low-luminosity counterpart within our sample, as has been seen in more-massive AGN.  We also find that the subset showing the presence of a blue wing in [\\ion{O}{3}] exhibits weak NLR emission.  We posit that in the objects with a blue wing, the outflow drives gas away from the central region, leading to a lower covering fraction \\cite[e.g.][]{2009ApJ...706..995L}.  However, contrary to our expectations from typical AGN, the presence of a blue wing does not correlate with high Eddington ratio, possibly due to our small dynamic range in this sample.  Lastly, we find that the high Eddington ratio composite has excess emission at intermediate velocities in the Balmer lines.  This corresponds to the ILR identified by \\cite{2008ApJ...683L.115H} and \\cite{2009ApJ...700.1173Z}.  This ILR emission could potentially bias $\\mbh$ estimates for AGN at high $L/\\ledd$ by decreasing the measured FWHM of the Balmer lines, thus underestimating $\\mbh$, as pointed out by \\cite{2006A&A...456...75C} and \\citet{2009ApJ...700.1173Z}.  Overall, the low-mass AGN studied here seem to behave in much the same way as more massive AGN, as traced by optical emission lines.  They are likely to have relatively low gas-phase metallicities in the NLR than their larger brethren, and may have steeper far-UV continuum slopes, but the structure and organization of their emission line regions seems to be largely similar."
    },
    "1208/1208.5926_arXiv.txt": {
        "abstract": "In this paper, we review the prospects for studies of active galactic nuclei (AGN) using the envisioned future Cherenkov Telescope Array (CTA). This review focuses on jetted AGN, which constitute the vast majority of AGN detected at gamma-ray energies. Future progress will be driven by the planned lower energy threshold for very high energy (VHE) gamma-ray detections to $\\sim 10$~GeV and improved flux sensitivity compared to current-generation Cherenkov Telescope facilities. We argue that CTA will enable substantial progress on gamma-ray population studies by deepening existing surveys both through increased flux sensitivity and by improving the chances of detecting a larger number of low-frequency peaked blazars because of the lower energy threshold. More detailed studies of the VHE gamma-ray spectral shape and variability might furthermore yield insight into unsolved questions concerning jet formation and composition, the acceleration of particles within relativistic jets, and the microphysics of the radiation mechanisms leading to the observable high-energy emission. The broad energy range covered by CTA includes energies where gamma-rays are unaffected from absorption while propagating in the extragalactic background light (EBL), and extends to an energy regime where VHE spectra are strongly distorted. This will help to reduce systematic effects in the spectra from different instruments, leading to a more reliable EBL determination, and hence will make it possible to constrain blazar models up to the highest energies with less ambiguity. ",
        "introduction": "\\label{Intro}  Active galactic nuclei (AGN) are extragalactic sources of enhanced activity that are powered by the release of gravitational energy from a supermassive central black hole. Energy linked to the black hole spin \\citep[e.g.,][]{bz77} or rotating accretion disks \\citep[e.g.,][]{bp82} may be instrumental for forming prominent jets which transport material from the innermost region of the AGN to kpc-, sometimes even Mpc-scale distances with relativistic speed. Such jets are usually identified through the detection of bright non-thermal radio emission as observed in radio-loud AGN.  Only a small percentage ($\\sim 10$~\\%) of all AGN are known to be radio-loud\\footnote{Radio-loud AGN are conventionally characterized with a radio-to-optical flux ratio $S_{\\rm{5GHz}}/S_{\\rm{440nm}}>10$.}. In the vicinity of the central region of an AGN matter is accreted from a disk onto the black hole, line-emitting clouds of material (the so-called broad-line region, BLR, and narrow line region, NLR) swirl at pc to kpc distances from the central engine, dusty material surrounding the accretion disk may imprint thermal signatures in the infrared part of the AGN spectrum, and the prominent jets of material in case of radio-loud AGN dominate the non-thermal radiative power in such systems (see Fig.~\\ref{upsketch}).  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{Fig1.eps} \\caption{Sketch illustrating the constituents and geometry of a radio-loud AGN \\citep[from][]{up95}.} \\label{upsketch} \\end{figure}  The radiation from material which moves relativistically with speed $\\beta_{\\Gamma} c$ (with $\\Gamma = 1/\\sqrt{1 - \\beta_{\\Gamma}^2}$ being the bulk Lorentz factor) along the jet axis is beamed into an angle $\\sim 1/\\Gamma$ around the direction of propagation. Because of this beaming effect mostly those AGN whose jet axes are close to alignment with our line of sight (i.e., blazars) are favourably detected as sources of high-energy (gamma-ray) emission. However, also some mis-aligned AGN (i.e., radio galaxies) can be detected, if they are sufficiently nearby. Blazars therefore offer an excellent opportunity to study jet physics of massive black hole systems, and through population studies also their evolution over cosmic time. Because the bolometric radiative energy output of AGN jets is often dominated by the gamma-ray regime\\footnote{We note that the apparent dominance of gamma rays in the overall blazar budget was recently shown to be affected by selection effects \\citep{giommi11}.}, the observed peak flux $(\\nu F_{\\nu})^{\\rm pk}$ in this band, together with a knowledge of the bulk Lorentz factor $\\Gamma$, provides a robust lower limit for the overall jet energetics constrained by the total radiative power, $L_{\\rm jet} > L_{\\rm rad}$, with \\begin{equation} L_{\\rm rad} \\approx {d_{\\rm{L}}^2 \\over \\Gamma^2} \\, (\\nu F_{\\nu})^{\\rm pk}, \\label{Lrad} \\end{equation} where $d_{\\rm{L}}$ is the luminosity distance. Such limit is not only crucial for constraining jet formation scenarios and the overall particle and field content of a jet including its impact for searches for the sources of the ultrahigh energy cosmic rays, but also for, e.g., investigating the jet's feedback on its environment. Comparing disk and jet energetics may give important clues on the physical connection between disk accretion and jetted outflows. Because these jets form in the vicinity of the strong gravitational fields of massive, probably rotating, black holes, studying events occuring close to the central engine may contribute to understanding jet formation. Size scales of the emission region of the order of the Schwarzschild radius are implied by extreme variability observed e.g., down to a few minutes time scales at TeV energies \\citep{aharonian07,albert07}, in a few radio-loud AGN, and this might imply a location of the emission region very close to the central black hole. On the other hand, the observation of systematic variations of the optical polarization over several days associated with a gamma-ray flare \\citep[e.g.,][]{3c279nature}, and distinct gamma-ray flares coinciding with the peak polarization of the mm-core \\citep{jorstad09} seem to favour rather pc-scale distances of the emission region relative to the central engine. This highlights the current debate regarding the location of the emission region. Studying the gamma rays from jets within the multifrequency context offers a view towards the global structure and composition of magnetized relativistic outflows, which provide constraints on the dominant radiation mechanisms. Monitoring the transition from flaring events to the quiescent phases together with the estimates on the overall flaring duty cycles may provide hints on the origin of variability.  Gamma rays probe the highest energy particles present in these jets, and therefore are relevant for our understanding of how charged particles are accelerated in jet plasmas, e.g., via shocks, and/or turbulence and/or magnetic reconnection. This may also have implications for our understanding of the origin of ultrahigh energy cosmic rays.  In this article, we review the prospects of CTA to facilitate progress in our understanding of the AGN phenomenon and its related physics including the large-scale impact of the associated jets. ",
        "conclusions": "Concluding remarks}   This surely incomplete list of topics discussed above reveals the potential of CTA for significant progress in the field of AGN research. Improvements in sensitivity and energy coverage will allow for the study of a much larger population of AGN, although we caution that the here important GeV energy range as is currently provided by the {\\it Fermi}-LAT instrument may be available at the time of CTA operations only to an extremely limited extent. This will enable to tackle a large range of topics from population studies and questions of cosmological evolution of AGN via studies of the formation and composition of extragalactic jets and the microphysics of the production of high energy emission in relativistic jets, to studies of the Extragalactic Background Light, which will shed light on the broader issues of cosmological galaxy evolution and structure formation. Most exciting, as CTA will enlarge the dynamical flux range and explore the high-redshift universe at VHEs, unexpected, possibly surprising, phenomena may challenge current theoretical concepts, and trigger to deepen our understanding of the extragalactic sky. This review might provide some insight into possible ways that observations by CTA --- coordinated with simultaneous observations at other wavelengths --- might lead to progress in the study of some of the most pressing questions of the VHE sky.  \\vspace*{1cm}  {\\bf{Acknowledgements}}  We like to thank Chuck Dermer, Benoit Lott, Marco Ajello and Paolo Giommi for providing excellent comments on this work which improved this manuscript. MB acknowledges support from NASA through Astrophysics Theory Program grant NNX10AC79G and Fermi Guest Investigator Grants NNX10AO49G and NNX11AO20G. AR acknowledges support by Marie Curie IRG grant 248037 within the FP7 Program."
    },
    "1208/1208.0732_arXiv.txt": {
        "abstract": "A leading formation scenario for R Coronae Borealis (RCB) stars invokes the merger of degenerate He and CO white dwarfs (WD) in a binary.  The observed ratio of $^{16}\\mathrm{O}$/$^{18}\\mathrm{O}$ for RCB stars is in the range of 0.3-20 much smaller than the solar value of $\\sim 500$.  In this paper, we investigate whether such a low ratio can be obtained in simulations of the merger of a CO and a He white dwarf.  We present the results of five 3-dimensional hydrodynamic simulations of the merger of a double white dwarf system where the total mass is $0.9 M_\\odot$ and the initial mass ratio (q) varies between 0.5 and 0.99.  We identify in simulations with $q\\lesssim0.7$ a feature around the merged stars where the temperatures and densities are suitable for forming $^{18}\\mathrm{O}$. However, more $^{16}\\mathrm{O}$ is being dredged-up from the C- and O-rich accretor during the merger than the amount of $^{18}\\mathrm{O}$ that is produced.  Therefore, on a dynamical time scale over which our hydrodynamics simulation runs, a $^{16}\\mathrm{O}/^{18}\\mathrm{O}$ ratio of $\\sim 2000$ in the ``best'' case is found.  If the conditions found in the hydrodynamic simulations persist for $10^6$ seconds the oxygen ratio drops to 16 in one case studied, while in a hundred years it drops to $\\sim 4$ in another case studied, consistent with the observed values in RCB stars. Therefore, the merger of two white dwarfs remains a strong candidate for the formation of these enigmatic stars. ",
        "introduction": "R Coronae Borealis stars (RCBs) are hydrogen deficient stars, with a carbon rich atmosphere \\citep{clayton96,clayton12}.  These very unusual stars are observed to be approximately $98\\%$ He and $1\\%$ C by mass.  The masses of RCB stars are difficult to measure since they have never been observed in a binary system, but stellar pulsation models have shown masses to be on the order of $1 M_\\odot$ \\citep{saio08,han98}.  The luminosity is characterized by a peculiar behavior: they fade at irregular intervals by up to 8 magnitudes, and gradually recover back to maximum luminosity over a period of a few months to a year. Such an observational feature is thought to be caused by clouds of carbon dust formed by the star itself \\citep{okeefe39}.  RCB stars show many anomalous elemental abundances compared to solar. Typically they are extremely deficient in hydrogen and are enriched relative to Fe, in N, Al, Na, Si, S, Ni, the s-process elements, and sometimes O \\citep{asplund00}.  The lower bound on the $^{12}$C/$^{13}$C ratio is between 14-100 for the majority of RCB stars, much larger than the equilibrium value in stars of solar metallicity which is 3.4 \\citep{hema12}, although at least one star, V CrA, shows a significant abundance of $^{13}$C \\citep{rao08,asplund00}.  Also, lithium has been detected in 5 RCB stars \\citep{asplund00,kipper06}. In other RCBs there is no lithium observed. The atmospheres of these stars show material processed during H burning via the CNO cycle and He burning via the 3-$\\alpha$ process.  In this paper, we focus on the more recent discovery of the oxygen isotopic ratio, $^{16}\\mathrm{O}$ to $^{18}\\mathrm{O}$ \\citep{clayton07,garcia09,garcia10}, found to be of order unity in RCB stars (the stars measured had ratios between 0.3 and 20). This ratio is found to be $\\sim$ 500 in the solar neighborhood \\citep{scott06}, and varies from 200 to 600 in the Galactic interstellar medium \\citep{wilson94}.  No other known class of stars displays $^{16}\\mathrm{O}/^{18}\\mathrm{O}\\sim$1 \\citep{clayton07}.  In a single star, partial He burning on the cool edge of the He burning shell  produces a significant amount of $^{18}$O but normally it would not be mixed to the surface. If He burning continues to its conclusion the $^{18}$O will be turned into $^{22}$Ne \\citep{clayton05}.  Two scenarios have been put forth to explain the progenitor evolution for RCB stars; one is a final helium shell flash and the other a double degenerate white dwarf (WD) merger \\citep{webbink84,renzini90}.  According to \\citet{iben96}, RCB stars could be the result of a final flash, when a single star late in its evolution has left the asymptotic giant branch and is cooling to form a WD and a shell of helium surrounding the core ignites.  However, the temperatures that result from He burning in a final flash will result in $^{14}$N being completely burned into $^{22}$Ne leaving little $^{18}\\mathrm{O}$ \\citep{clayton07}.  In the second scenario, a close binary system consisting of a He and a CO WD merges, leading to an RCB star \\citep{webbink84}. \\citet{iben96} explain that theoretically the accretion of a He WD $\\sim 0.3 M_{\\odot}$ onto a CO WD $\\sim0.6 M_{\\odot}$ can produce a carbon-rich supergiant star (M$\\sim$0.9 M$_{\\sun}$) that is hydrogen-deficient at the surface after two common envelope (CE) phases.  In this scenario, the He WD is disrupted and forms the envelope of the newly merged star while the CO WD forms the core.  In He burning conditions the partial completion of the reaction chain ($^{14}\\mathrm{N(\\alpha,\\gamma)}^{18}\\mathrm{F(\\beta^+)} ^{18}\\mathrm{O(\\alpha,\\gamma)}^{22}\\mathrm{Ne}$) will take place during the merger when accretion results in high temperatures, and C and O may be dredged up from the core.  A large amount of $^{18}\\mathrm{O}$ will be created only if this process is transient and is not allowed to proceed to completion.  The available $^{14}\\mathrm{N}$ is a result of CNO cycling in the progenitor star, and the amount depends on the initial metallicity of that star.  Hence the maximum amount of $^{18}\\mathrm{O}$ formed cannot exceed the initial abundance of $^{14}\\mathrm{N}$, unless additional $^{14}\\mathrm{N}$ can be produced. Our objective is to investigate whether the merger of a He WD and a CO WD with a combined mass of $0.9M_\\odot$ \\citep[similar to RCB star masses;][]{saio08} can lead to conditions suitable for producing oxygen isotopic ratios observed in RCB stars.  Close WD binary systems may be the progenitors for type Ia supernovae \\citep{iben84}, and such systems have therefore attracted much interest both from theoretical and observational points of view.  Using an earlier version of the hydrodynamics code used in this work, \\citet{motl07} and \\citet{dsouza06} studied the stability of the mass transfer in close WD binary systems.  The fate of a close WD binary system depends on the mass ratio of the two WDs.  Neglecting the angular momentum in the spin of the binary components and allowing the angular momentum contained in the mass transfer stream to be returned to the orbit, \\citet{paczynski67} found that mass ratios below $2/3$ are stable.  However, if the mass transfer stream directly strikes the accretor instead of orbiting around it to form an accretion disk, this stability limit may be reduced significantly.  Recent simulations by Marcello et al.  (private comm.) indicate that even a mass ratio of 0.4 may be unstable and lead to a merger. \\citet{brown11} have recently reported observations of the close WD binary system, SDSS J065133.33+284423.3, which consists of a $0.25 M_\\odot$ He WD and a $0.55 M_\\odot$ CO WD, and is predicted to start mass transfer in about 900,000 years.  If these two WDs merge, what will the resulting object look like?  Other groups have employed smooth particle hydrodynamics (SPH) simulations to study WD mergers, for instance \\citet{benz90}, \\citet{yoon07}, and \\citet{raskin11}. In \\citet{motl12}, the results from grid based hydrodynamics simulations are compared to SPH simulations of WD mergers, and it is found that the two methods produce results in excellent agreement. Very recently, \\citet{longland11} studied the nucleosynthesis as the result of the merger of a $0.8 M_\\odot$ CO WD and a $0.4 M_\\odot$ He WD.  The merger simulation was performed with an SPH simulation code \\citep{loren-aguilar09}.  They find that if only the outer part of the envelope ($0.014<R/R_\\odot<0.05$) is convective, the $^{16}\\mathrm{O}$ to $^{18}\\mathrm{O}$ ratio is 19, which is in the range measured for RCB stars \\citep{clayton07}.  On the contrary, if the entire envelope is convective, the ratio is 370.  \\citet{jeffery11} investigated the surface elements resulting from a merger of a CO with a He WD, based on 1-D stellar evolution models and parametric nucleosynthesis analysis.  They considered two situations, a cold (no nucleosynthesis) merger, and a hot merger (with nucleosynthesis).  In both cases, they find surface abundances of C, N, and O that can be made to match the observed RCB star surface abundances.  S and Si however, do not match\\footnote{\\citet{asplund00} sugests that a possible solution is condensation of dust which removes some gas phase abundance.}. In the hot merger scenario, the most promising location for nucleosynthesis to take place is in a hot and dense region just on the outside of the original accretor, as for instance seen in the simulations by \\citet{yoon07}, \\citet{loren-aguilar09}, or \\citet{raskin11}.  This region forms as accreting matter from the donor impacts the accretor.   In this paper, we investigate whether the unusual abundances measured in RCB stars can be produced in a WD merger, by first performing hydrodynamic simulations of the merger of two WDs \\citep[using a a modified version of the 3-dimensional hydrodynamic code called Flower, see][]{motl02}.  In section 2 of this paper, we present the methodology of our work.  Here the details of the hydrodynamic code and the nucleosynthesis code along with their initial conditions are given.  In section 3, the results of the hydrodynamic simulations and their corresponding nucleosynthesis calculations are presented.  Finally, in section 4 we compare our results with the results of other authors, and in section 5 we discuss our results along with future directions for the work in this paper. ",
        "conclusions": "The ratio of $^{16}{\\mathrm O}$ to $^{18}{\\mathrm O}$ is observed to be very low (of the order unity) in RCB stars.  From hydrodynamic simulations of double degenerate mergers for various mass ratios, we have therefore looked for conditions that would allow for production of $^{18}{\\mathrm O}$ in order to explain this ratio.  This requires temperatures of the order $\\sim1.5$ - $3\\times10^8$ K.  At lower temperatures, $^{18}{\\mathrm O}$ will not form on the available dynamic time scale, and at higher temperatures it will be destroyed (converted to $^{22}\\mathrm{Ne}$) as soon as it forms.  Our results show that the maximum temperature that can be found in the SOF depends strongly on the mass ratio, $q$.  Low values of $q$ give temperatures in the SOF of $1.2-2.5 \\times10^8$ K (assuming realistic abundances), while mass ratios above $q=0.7$ gives temperatures much lower than that.  Hence the lower $q$ values that we have investigated will give temperatures in the SOF suitable for $^{18}{\\mathrm O}$ production on a dynamical timescale.  On a dynamical time scale, we do not find very low oxygen ratios in the SOF in any of our simulations.  After a thousand seconds, the $q=0.5$ simulation reaches a $^{16}\\mathrm{O}$ to $^{18}\\mathrm{O}$ ratio of about 2000, since not much $^{18}\\mathrm{O}$ is produced on such a short time scale and because of the large amount of $^{16}\\mathrm{O}$ present from the dredge up.  Within a day, the oxygen ratio in the $q=0.5$ simulation reaches its lowest value of $\\sim30$ (Fig.~\\ref{img_o16_o18}). After that, $^{18}\\mathrm{O}$ is being destroyed and the ratio increases.  The $q=0.6$ simulation reaches its lowest value of 16 after $10^6-10^7$ seconds, after which the ratio increases again as $^{18}\\mathrm{O}$ is being destroyed.  The best case for obtaining low oxygen ratios is the $q=0.7$ simulation, which reaches the lowest $^{16}\\mathrm{O}$ to $^{18}\\mathrm{O}$ of $4$ after $\\sim10^{2}$ years, assuming the conditions remain constant for that long.  This is comparable to the observed oxygen ratios in RCB stars. As we have shown, the nuclear reactions begin very fast, and the extra energy released is likely to affect the conditions in the SOF. Unfortunately we cannot model this at present, and here we simply state that if the very low oxygen ratios of order unity shall be achieved in the $q=0.7$ simulations, the conditions in the SOF must remain relatively unchanged for a period of about a hundred years.  In the high-$q$ simulations the temperature was not sufficiently high to produce  $^{18}\\mathrm{O}$. However, as we have shown the protons react very quickly, releasing an amount of energy comparable to the thermal energy in the SOF. Hence, it is possible that even in the high $q$ simulations, the SOF prior to the merger can become sufficiently hot for the nucleosynthesis involving helium to start.  This extra energy from nuclear processes leads to an extra pressure term, that could potentially help preserve the SOF also in the high-$q$ simulations so that the $^{18}\\mathrm{O}$ production can continue. Furthermore, even though the amount of $^{18}\\mathrm{O}$ in RCBs is strongly enhanced compared to all other known objects, it is important to keep in mind that among RCB stars the observed oxygen ratio varies by 2 orders of magnitude from star to star.  The $^{16}{\\mathrm O}$ to $^{18}{\\mathrm O}$ ratio depends both on the formation of $^{18}{\\mathrm O}$, and also on the amount of $^{16}{\\mathrm O}$ present.  In all of our simulations, we have found significant dredge up of accretor material which consists primarily of $^{16}{\\mathrm O}$ and $^{12}{\\mathrm C}$. $^{18}{\\mathrm O}$ is formed from  $^{14}{\\mathrm N}$, which in part depends on the initial metallicity of the progenitor stars, assumed to be solar giving about $1\\%$ of $^{14}{\\mathrm N}$. But $^{14}\\mathrm{N}$ is also being produced (Figs.~\\ref{img_0.5}, \\ref{img_0.6}, and \\ref{img_0.7}) in the nuclear processes occurring in the SOF. The $^{16}\\mathrm{O}$ to $^{18}\\mathrm{O}$ ratios, that we find in the SOF (shown in Figs.~\\ref{img_o16_o18}),  are the lowest values possible from our simulations.  If less $^{16}\\mathrm{O}$ were dredged up from the accretor, then less $^{18}\\mathrm{O}$ would need to be produced in order to obtain low oxygen ratios.  One way to avoid contaminating the SOF with $^{16}\\mathrm{O}$ from dredge-up, would be if the accretor is a hybrid He/CO WD.  \\citet{rappaport09} modeled a $0.475 M_\\odot$ hybrid He/CO WD with a He envelope of more than $0.1 M_\\odot$.  In our simulations, we found that about $0.1 M_\\odot$ of accretor material was dredged-up, most of it ending up in the SOF.  If the accretor is a hybrid He/CO WD, most of this dredged-up material might turn out to be $^4\\mathrm{He}$, and the $^{16}\\mathrm{O}$ contamination of the SOF would be much less.  In this case, the lower mass of the accretor means the donor must also have a lower mass in order to get a tidal disruption of the donor rather than a core merging.  This lower total mass could result in lower temperatures.  Furthermore, a mixture with mostly He could also lead to lower temperatures (Eqs.~\\ref{eq:T} and \\ref{eq:cv}).  However, as we discussed above, the reacting protons might heat the gas to a sufficiently high temperature to start the helium burning no matter what the total mass. Hence, it remains to be seen if the oxygen isotope ratio will be of the correct order if this is the situation. \\citet{han02, han03} showed that a significant fraction of sdB stars are expected to be in close-period binaries with He WDs.  After completion of helium burning, sdB stars may evolve into hybrid CO/He WDs \\citep{justham11}. In \\citet{clausen12}, a population synthesis study found that there should be almost 200,000 sdB stars in binaries in the galaxy.  Hence it may not be unreasonable that a small fraction of these are in a close binary with a He WD, and that after the sdB star has evolved into a hybrid WD they can merge.    When the merged object begins to expand, the temperature and density of the SOF are likely to drop which may bring the nuclear processes there to a halt. Hence the oxygen ratio in the SOF at that time may be representative of what will be observed at the surface of this object. With our numerical approach, we cannot determine when the merged object will begin to expand.  Convection may be triggered quickly in the SOF, as the thermal diffusion time \\citep[$\\tau_{\\rm th}\\sim10^9$ seconds using $\\varepsilon_{\\rm th}\\sim 10^7 {\\rm ~erg~g^{-1}~s^{-1}}$ from][]{yoon07} is much larger than the 100 seconds thermonuclear time scale ($\\tau_{\\rm nuc}=c_p T/\\epsilon_{\\rm nuc}$, where $c_p$ is the specific heat at constant pressure that we estimate to be of order $10^8$ erg/g K, $T$ is the temperature of order $10^8$ K, and $\\epsilon_{\\rm nuc}$ is the nuclear energy production rate of order $10^{14} {\\rm ~erg~g^{-1}~s^{-1}}$ after 1000 seconds; Fig.\\ref{Energy}).  If the material outside the merged core and the SOF accretes quickly, it could affect the SOF.  \\citet{ken12} argued that the viscous time scale for material in a disk surrounding such a core is an hour to a year, indicating that it could accrete relatively fast. Alternatively, the merged object can expand before this material has accreted.  In that case, the SOF is unlikely to be affected much by this material.  However, the \\citet{ken12} result would indicate that the expansion would have to begin quickly (less than a year) since the accretion may occur on this time scale. We found in the $q=0.6$ simulation that the oxygen ratio is 16 after only $10^6$ seconds, a significant enhancement from the solar value and comparable to what is seen in some RCBs.  We note that most of the accretor material, that is being dredged up, ends up in the SOF.  The material outside the SOF is mostly from the He WD (Table~\\ref{dredgeuptable}).  The conditions outside the SOF are not favorable for nuclear processes to occur, and so if only the matter outside the SOF swells up to supergiant size (without mixing in newly produced elements in the SOF), the resulting object could look like a helium star with some carbon, ie. an RCB star.  We have found that the best conditions for reaching the low oxygen ratios occur if the temperature remains $\\sim10^8$ K.  Then $^{18}\\mathrm{O}$ forms slowly, and is destroyed even more slowly.  Hence it is plausible that the merger of a CO WD and a He WD can lead to the formation of an RCB star, although a more elaborate investigation must be performed in order to conclusively answer this question."
    },
    "1208/1208.5051_arXiv.txt": {
        "abstract": "We report the detection of orbital decay in the 12.75-min, detached binary white dwarf (WD) SDSS J065133.338+284423.37 (hereafter J0651). Our photometric observations over a 13-month baseline constrain the orbital period to $765.206543(55)$ s and indicate the orbit is decreasing at a rate of $(-9.8\\pm2.8) \\times 10^{-12}$ s s$^{-1}$ (or $-0.31\\pm0.09$ ms yr$^{-1}$). We revise the system parameters based on our new photometric and spectroscopic observations: J0651 contains two WDs with $M_1 = 0.26\\pm0.04$ \\msun\\ and $M_2 = 0.50\\pm0.04$ \\msun. General relativity predicts orbital decay due to gravitational wave radiation of $(-8.2\\pm1.7) \\times 10^{-12}$ s s$^{-1}$ (or $-0.26\\pm0.05$ ms yr$^{-1}$). Our observed rate of orbital decay is consistent with this expectation. J0651 is currently the second-loudest gravitational wave source known in the milli-Hertz range and the loudest non-interacting binary, which makes it an excellent verification source for future missions aimed at directly detecting gravitational waves. Our work establishes the feasibility of monitoring this system's orbital period decay at optical wavelengths. ",
        "introduction": "The 12.75-minute orbital period detached binary WD system J0651 was discovered by \\citet{BrownJ0651} as part of the extremely low mass (ELM, $\\leq 0.25$ \\msun) WD Survey, a targeted spectroscopic search for ELM WDs. While that survey has yielded some two dozen merger systems, with orbital periods of tens of minutes to hours \\citep{BrownELMi,KilicELMii,BrownELMiii,KilicELMiv}, none are as compact as J0651.  In addition to a large radial-velocity amplitude, this double degenerate system is oriented in such a way that it yields a wealth of photometric information: eclipses of each star by the other, ellipsoidal variations and Doppler boosting. While photometric observations engender an accurate way to measure the orbital and system parameters, they also provide multiple clocks with which to monitor the orbital evolution of the system.  The orbital decay of compact binary systems is currently the best method to detect the influence (and existence) of gravitational waves, and few known systems are radiating as strongly or decaying as rapidly as J0651. There are presently just five binaries known with orbital periods less than 15 minutes, and the other four are interacting: three are the AM CVn systems HM Cnc, V407 Vul, and ES Cet, and the other is the low-mass X-ray binary 4U 1820-30 \\citep{Israel99,Haberl95,Steeghs06,Warner02,Stella87}. After the 5.4-minute HM Cnc \\citep{Israel02,Roelofs10}, J0651 is the second-loudest gravitational wave source known in the milli-Hertz frequency range \\citep{Amaro12}. J0651 is thus the shortest-period detached compact binary known and the cleanest system to observe at optical wavelengths for orbital decay due to gravitational wave radiation.  In this Letter we present follow-up photometric and spectroscopic observations of J0651, refine orbital and system parameters, and report the detection of rapid orbital decay in the system. We discuss the orbital period change in the context of expectations from general relativity, as well as deviations expected due to tidal interactions. Section 2 describes our observations, and Sections 3 and 4 present the refined system parameters and the orbital period decay, respectively. ",
        "conclusions": "Based on the refined parameters for the J0651 system (see Section~\\ref{sec:param}) and treating each WD as a point mass in a non-relativistic circular orbit, general relativity predicts an orbital period decay in this system of $(-8.2\\pm1.7) \\times 10^{-12}$ s s$^{-1}$ (or $-0.26\\pm0.05$ ms yr$^{-1}$) \\citep{LL75}. Recently, \\citet{vdB12} demonstrated that the point mass approximation is valid to better than 1\\% in cases such as J0651, so our uncertainty in the orbital decay from gravitational wave radiation is dominated by the uncertainty in the component masses. Additional effects that could modulate the mid-eclipse times are unlikely to explain the observed shift: For example, given the estimated distance of 1.0 kpc \\citep{BrownJ0651}, we expect the proper motion to change the period by no more than $5 \\times 10^{-16}$ s s$^{-1}$ \\citep{Shklovskii70}.  It is evident from the high-amplitude ellipsoidal variations of the primary that strong tidal forces are also present. These tides will act as a torque to spin-up the WDs if the system is synchronized, further robbing the orbit of angular momentum and increasing the rate of orbital period decay. The degree to which this tidal torquing influences the orbital evolution depends on the effective tidal locking, which is in many ways determined by the physical structure of the ELM WD. This effect could increase the rate of period decay by at least 5\\% if the system is synchronized \\citep{Piro11,Benacquista11,Fuller12}. With just 13 months of monitoring, our sensitivity in the observed rate of orbital decay is not yet sufficient to detect a significant deviation from pure gravitational wave losses. However, future observations should constrain this discrepancy, providing an excellent probe of the interior of ELM WDs, in addition to the possibilities opened through asteroseismology \\citep{Steinfadt10,Hermes12}.  The short period of J0651 makes it one of the loudest known sources of gravitational wave radiation, and continued monitoring of orbital decay in the system will provide strong constraints on the gravitational wave strain of J0651. This is important for future gravitational wave missions, like the evolved Laser Interferometry Space Antenna ({\\em eLISA}). Critical to that success is disentangling the contributions of tidal torques on the orbital decay, an effort worthy of further photometric observations and modeling."
    },
    "1208/1208.5267_arXiv.txt": {
        "abstract": "{ The origin of coronal type-II radio bursts and the nature of their band-splitting are still not fully understood, though a number of scenarios were proposed to explain them. This is largely due to the lack of detailed spatially resolved observations of type-II burst sources and of their relations to magnetoplasma structure dynamics in parental active regions. } { To make progress in solving this problem on the basis of one extremely well observed solar eruptive event. } { The relative dynamics of multi-thermal eruptive plasmas, observed in detail by the Atmospheric Imaging Assembly onboard the Solar Dynamics Observatory and of the harmonic type-II burst sources, observed by the Nan{\\c c}ay Radioheliograph at ten frequencies from 445 to 151~MHz, is studied for the 3 November 2010 event arising from an active region behind the east solar limb. Special attention is given to the band-splitting of the burst. Analysis is supplemented by investigation of coronal hard X-ray (HXR) sources observed by the Reuven Ramaty High-Energy Solar Spectroscopic Imager. } { It is found that the flare impulsive phase was accompanied by the formation of a double coronal HXR source, whose upper part coincided with the hot ($T \\approx 10$~MK) eruptive plasma blob. The leading edge (LE) of the eruptive plasmas ($T \\approx 1-2$~MK) moved upward from the flare region with the speed of $v \\approx 900-1400$~km~s$^{-1}$. The type II burst source initially appeared just above the LE apex and moved with the same speed and in the same direction. After $\\approx20$~s it started to move about twice faster, but still in the same direction. At any given moment the low frequency component (LFC) source of the splitted type-II burst was situated above the high frequency component (HFC) source, which in turn was situated above the LE. It is also found that at a given frequency the HFC source was located slightly closer to the photosphere than the LFC source.  } { Based on the set of established observational facts it is concluded that the shock wave, which could be responsible for the observed type-II radio burst, was initially driven by the multi-temperature eruptive plasmas, but later transformed to a freely propagating blast shock wave. The most preferable interpretation of the type-II burst splitting is that its LFC was emitted from the upstream region of the shock, whereas the HFC -- from the downstream region. The shock wave in this case could be subcritical. } ",
        "introduction": "\\label{Introduction}  It is generally accepted that coronal type-II radio bursts are a signature of MHD shock waves \\cite[\\eg{},][]{Zheleznyakov70,Wild72,Nelson85,Mann95,Cairns11}. Nevertheless, there is a long-lasting question, whether these shock waves are 1) blast shocks due to explosive flare energy releases or 2) piston shocks driven by eruptive magnetoplasma structures often evolving into coronal mass ejections (CMEs). Numerous observational evidences were reported in favor of both the first and the second scenario \\citep{Nelson85,Aurass97,Cliver99,Vrsnak08,Vourlidas09}. To make progress in solving this problem, spatially resolved observations of type-II burst sources at several frequencies in the lower corona ($\\lesssim 2R_{\\odot}$) are required. Radio observations must be supplemented by high-precision, high-cadence, multi-wavelength observations of magnetoplasma structure dynamics in parental active regions. It is obvious that limb events are best suited for this purpose.  The partially behind the East limb solar flare of 3 November 2010 (C4.9~class, $\\approx$12:10~UT) is a prominent candidate to satisfy these requirements. It was accompanied by a split-band decimetric/metric type-II radio burst, whose sources were well observed by the Nan{\\c c}ay Radioheliograph \\citep[NRH]{Kerdraon97} in all ten working frequencies from 445 up to $\\approx$151 MHz. An important fact is that the radio sources were observed at low altitudes between $\\approx$0.25 R$_\\odot$ and $\\approx$0.65 R$_\\odot$. Moreover, the Atmospheric Imaging Assembly \\citep[AIA;][]{Lemen11} onboard the Solar Dynamics Observatory observed an erupting multi-thermal magnetoplasma structure in the lower corona $\\lesssim$0.4 R$_\\odot$ \\citep[see][for details]{Reeves11,Foullon11,Cheng11}.  The nature of the band-splitting effect, often observed in coronal and interplanetary type-II bursts, is still an unsolved riddle, although several mechanisms were proposed to explain it \\citep[see, \\eg{},][as reviews]{Kruger79,Nelson85,Vrsnak01,Cairns11}. The first popular mechanism, initially proposed by \\citet{McLean67}, assumes the existence of two (or more) regions with different physical characteristics (say, plasma concentration) along a shock front. The second popular mechanism was proposed by \\citet{Smerd74,Smerd75}. It suggests that the two sub-bands of a splitted coronal type-II burst are due to coherent plasma radio emission simultaneously generated in the upstream and downstream regions of a shock wave. The event of 2010 November 3 gives us an opportunity to investigate this interesting and important effect, which in turn can be closely related to the entire problem of the type-II bursts origin.  It should be noted that combined analysis of this type-II radio burst on the 3rd of November 2010 and associated plasma eruption was reported by \\cite{Bain12} very recently. However, \\citeauthor{Bain12} mainly concentrated on the dynamics of the most intense type-II burst sources. Our study is more focused on the band-splitting effect.  The paper is organized as follows. Analysis of the observational data is described in Section~\\ref{Observations}. Results of the analysis are summarized in Section~\\ref{Results}. These results are discussed and interpreted in Section~\\ref{Discussion}. Final remarks are given in Section~\\ref{FR}. ",
        "conclusions": "\\label{Discussion}  First of all, it should be noted that according to the AIA/SDO and RHESSI combined observations, the entire picture of the 3 November 2010 event was consistent with the standard eruptive flare scenario, \\ie{}, the CSHKP model \\citep[see also][]{Reeves11,Foullon11,Cheng11}. Shock waves of various kinds are expected phenomena of such a model \\citep[\\eg{},][]{Hirayama74,Priest82}.  However, we shall also clarify here that there are no direct observational evidences of a shock in this event. No shock wave fronts which would have sharply separated the ``disturbed'' (downstream) and ``undisturbed'' (upstream) regions ahead of the leading edge of the erupting plasmas were found using EUV observations of the AIA/SDO. We emphasize that the measured leading edge of the eruptive plasmas is most probably not a hypothetical shock wave front. If a shock wave was indeed formed, its front should be located at each time somewhere above the measured leading edge. Nevertheless, a compelling indirect evidence of a shock wave formation was found in this event with the NRH and Phoenix observations --- the type-II radio burst sources were observed to propagate with (probably) super-magnetosonic speeds above the leading egde of the eruptive plasmas. We will thus assume below that a shock wave was really formed in this event.  Our discussion below will be limited to the following two subjects: 1) possible origin of the shock wave, 2) possible origin of the observed band-splitting of the type-II radio burst.  \\subsection{Shock wave driver} \\label{Driver}  As it has already been mentioned in Section~\\ref{Introduction}, type-II radio bursts are believed to be produced by MHD shock waves. However, it is still unclear whether these shock waves are 1) bow or piston shocks driven by eruptive magnetoplasma structures or 2) blast shocks due to an explosive flare energy release localised in space and time \\citep[see comprehensive reviews on this topic given by, \\eg{},][]{Vrsnak00,Vrsnak08}.  \\subsubsection{Piston-driven shock wave} \\label{PDSW}  We found strong evidence in favour of the first (piston-driven shock wave) scenario: 1) the location of the type-II burst sources above the apex of the eruptive plasma leading edge (see  Figure~\\ref{fig6}, Figure~\\ref{fig7}, Figure~\\ref{fig8}), and 2) the same propagation direction of the type-II burst sources and of the erupting plasma apex (see  Figure~\\ref{fig3}, Figure~\\ref{fig5}, Figure~\\ref{fig6}).  These results are similar to the results obtained by \\citet{Bain12} for the same event of 3 November 2010 and by \\citet{Dauphin06} for the 3 November 2003 flare, when the type-II burst sources were observed originating above the rising soft X-ray loops. Both \\citet{Bain12} and \\citet{Dauphin06} interpreted their observations in the frame of the piston-driven shock wave scenario. Spatially resolved observations of the coronal type-II burst sources in a close association with the propagating disturbances observed in the soft X-ray range were also reported earlier by, \\eg{}, \\citet{Klein99,Khan02}, and it was argued that the shock waves, which could be responsible for the type-II bursts, were most probably driven by these disturbances.  The fact that the estimated speed of the agent which excited the low frequency component source (LFC) of the type-II burst was much larger than the speed of the supposed driver ($v_{LFC} \\approx 2v_{LE} \\approx 2200$~km~s$^{-1}$), \\ie{} the erupting plasma, is not contradictory to the piston-driven scenario. The shock front velocity is indeed expected to be equal to the driver velocity only in the case when the driver has a constant geometrical shape, when it propagates in a homogeneous background medium, and also when the oncoming background plasma can wrap the driver's body. A bow shock wave is formed in this case \\citep[\\eg{},][]{Vrsnak00}. However, in the present event the possible shock driver is propagating in the stratified solar atmosphere with decreasing magnetic field and plasma density. Its geometrical size increases in time (see Figure~\\ref{fig6}), and the upstream background plasma can not easily pass the eruptive warm plasma rim because the later seems to remain rooted to the photosphere by magnetic field lines. Thus, the driver looks like a 3D expanding piston (probably not a spherical one but rather flux-rope shaped) which also has a preferred (upward) direction of motion. A situation when a shock wave propagates through the gravitationally stratified corona including dense loops was numerically simulated, \\eg{}, by \\citet{Pohjolainen08,Pomoell08,Pomoell09}. It was found that under these conditions the shock wave can propagate faster through the corona than its driver (see Subsection~\\ref{S1} for further discussion of this issue). It was also found that the shock front is strongest near the leading edge of the erupting plasma.   \\subsubsection{Blast shock wave} \\label{BSW}  Contrary to the piston-driven shock wave scenario, only one piece of indirect evidence (see Result~11 in Section~\\ref{Results}) could support the second (blast shock wave) scenario. It implies that a free-propagating blast shock wave could in principle be generated by the impulsive flare energy release which was evidenced by the formation of the coronal hard X-ray sources in the flare impulsive phase. Note, that \\citet{Bain12} did not completely rule out this possibility for the present event of 3 November 2010. \\citet{Vrsnak06} also reported evidence in favour of the blast shock wave scenario for the 3 November 2003 similar event. On the other hand, Result~11 may be just a coincidence. It alone seems not enough to interpret the studied type-II burst in frame of the blast shock wave hypothesis. Moreover, it is obvious that the blast shock wave hypothesis can not easily explain the same propagation direction of the type-II burst sources and of the erupting plasma apex. For these reasons we conclude that the type-II burst studied here was most probably produced, at least in the initial stage, by a piston-driven shock wave rather than by a blast shock wave.  Additional comments on the assumption of the piston-driven scenario for this event are found below. The value of $v_{LFC}\\approx2200$~km~s$^{-1}$ was obtained using the linear least square approximation of all ten observational data points for the LFC sources (shown by green marks on the height-time plot (Figure~\\ref{fig8})). However, the first (in time) half of these data points in Figure~\\ref{fig8} seems to behave differently than the second half, thus indicating that a linear approximation is probably not the best one. If the linear approximation is performed only on the first group of data points, the slope of this linear fit line is much closer (although a little bit steeper) to the slope of the linear fit for the leading edge of the eruptive plasma (red and black lines on Figure 8). This means that during the beginning of the type-II burst the characteristic velocity of its sources was similar to the velocity of the eruptive plasma leading edge. About 20~s later it became about twice the magnitude, \\ie{} $v_{LFC}\\approx2200$~km~s$^{-1}$.  Recent numerical simulations of \\citet[][]{Pomoell08,Pomoell09} have shown that the shock wave velocity can quickly exceed the velocity of the driver (flux rope) and then the shock escapes from the driver. In other words, it was found that the shock wave can be of the piston-driven type in the beginning of eruption and after some time it can transform to the freely propagating blast wave. These findings of \\citet[][]{Pomoell08,Pomoell09} are very similar to our observations.  It should also be mentioned here that the discussed change of slope in the type II data points occurred at the time when the band-splitting was first observed. It is possible that these two facts could be related to each other.     \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{fig9_bw_rgb.eps}} \\caption{ Schematic illustration of the 3 November 2010 eruptive event observations combined with their interpretation in the frame of the upstream-downstream scenario (see text). View is from the Heliographic North Pole. Direction to the Earth is marked by a thick black arrow. Notations: (1) hypothetical shock wave, (2) LFC source of the type-II burst, (3) its HFC source, (4) turbulent magnetosheath, (5) warm ($T \\approx 1-2$~MK) plasma rim and (6) its leading edge, (7) hot ($T \\simeq 10$~MK) erupting flux rope or plasma blob if it is observed from the Earth, (8) photosphere. Thin black arrows show directions of the eruptive plasmas, shock wave, LFC and HFC sources motion. Lengths of the arrows are proportional to the corresponding velocities of motion. Levels of constant undisturbed background electron plasma concentration, assuming the natural gravitational stratification, are marked by black dashed arc-lines, and $n_{1}>n_{2}>n_{3}$. } \\label{fig9} \\end{figure}   \\subsection{Split-band effect} \\label{Splitting}  It is not clear yet, which physical mechanism is responsible for the splitting of type-II bursts' emission. Currently, two main interpretations dominate. This does not mean, of course, that any alternative idea may not be valid \\citep[\\eg{},][and references therein]{Treumann92,Cairns11}.  One popular interpretation (henceforth Scenario~1) was proposed by \\citet{Smerd74,Smerd75}. It was suggested that the two sub-bands of splitted coronal type-II bursts, the LFC and HFC according to our terminology, could be due to coherent plasma radio emission simultaneously generated ahead of and behind a shock wave front, \\ie{}, in the upstream and downstream regions, respectively.  Another popular interpretation (henceforth Scenario~2), initially proposed by \\citet{McLean67}, suggests that different parts of a shock wave front could simultaneously encounter coronal structures of different physical properties, such as electron plasma density or magnetic field. In the case when different parts of a shock wave front propagate through the corona in more than two media with different physical conditions, it is expected to observe a type-II burst with a multiple set of bands. Different situations could occur. For example, those parts of the shock front which are parallel to surfaces of constant electron density would emit more intensively in some narrow frequency ranges than in others. In particular, \\citet{McLean67} simulated an idealized situation of a shock front encounter with a streamer and could reproduce split-band of type-II bursts. Similar ideas based on the shock drift acceleration mechanism have been discussed by, \\eg{}, \\citet{Holman83}. Recently, more sophisticated but ideologically similar numerical experiments of \\citet{Knock05} also reproduced splitting of coronal type-II bursts.    \\subsubsection{Scenario~1} \\label{S1}  It seems that our observations geometrically support Scenario~1 more than Scenario~2. The major argument in favour of the upstream-downstream scenario of \\citet{Smerd74,Smerd75} is that at every time the low frequency component (LFC) source is  located above the high frequency component (HFC) one, and that the HFC source fills almost all the space between the LFC source and the leading edge of eruptive plasmas (see Figure~\\ref{fig7} and Figure~\\ref{fig8})\\footnote{It should be noted here that the HFC source at 327.0~MHz appears to be located inside the warm eruptive plasma rim in Figure~\\ref{fig7}(b). This could be due to projection of the curved HFC source (because of a curved hypothetical shock wave front) into the image plane. See Figure~\\ref{fig9} as an illustration.}. This is schematically illustrated in Figure~\\ref{fig9}. In this case the LFC source, situated in the upstream region, can be naturally explained in the frame of some standard shock wave theories, \\ie{}, the shock drift acceleration mechanism.  There is another important argument in favour of Scenario~1. It was found (see Subsection~\\ref{T2B} and Figure~\\ref{fig5}) that at a given frequency the average position of the HFC source centroid was a little bit closer to the photosphere (and the flare site) than the average position of the LFC source at the same frequency. Although this is a subtle effect, it shows that at any time the plasma density in the region, from where the HFC sources are emitted, is enhanced relative to the undisturbed background plasma density. In our opinion, the natural explanation of this effect is that at a given frequency the HFC sources were emitted below the shock wave front --- in the downstream region, \\ie{}, in the magnetosheath, whereas the LFC sources were emitted at the same frequency in the upstream region. This idea is schematically illustrated in Figure~\\ref{fig10}.  It can be recalled here that both similar and opposite behaviours of the LFC and HFC sources were reported earlier in the literature \\citep[\\eg{},][]{Dulk70,Nelson75,Aurass97,Khan02}. In some cases the LFC sources are located farther from the flare site (or from the photosphere) than the HFC sources at the same frequencies but in some other cases they are closer to or almost at the same positions. All these earlier observations (known to us) were made at frequencies below $\\approx160$~MHz or for flare regions which were located close to the center of the visible solar disk. This makes it difficult to carry out a direct analogy between these and our observations.   \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{fig10_bw_rgb.eps}} \\caption{ Schematic illustration of the lower location of the HFC sources of the type-II burst with respect to the LFC sources observed at the same frequency. Panel (a) corresponds to an instant $t_{1}$, which is before the instant $t_{2}$ ($t_{2}>t_{1}$) of panel (b). Solar surface is depicted by thick black line with inclined ticks on the right. Thick dashed arc-like line shows the shock wave front. Light and dark grey ellipses represent the LFC and HFC sources, respectively. Horizontal arrow shows the direction of their movement. Levels of constant plasma density are shown by vertical straight dashed and dotted lines. Corresponding plasma densities are marked by $n_{1}$, $n_{1}'$, $n_{2}$, with $n_{1}'>n_{1}$ and $n_{1}>n_{2}$. Corresponding second harmonic of plasma frequencies, at which the LFC and HFC sources are emitted, are marked by $f_{1}$ and $f_{2}$, with $f_{1}>f_{2}$. } \\label{fig10} \\end{figure}    Scenario~1 meets however with a couple of difficulties. The generation of HFC emission requires intense electron plasma waves in the downstream region, whereas in situ measurements near the interplanetary shocks reveal them mainly in the upstream region \\citep[\\eg{},][]{Bale99,Thejappa00,Hoang07,Pulupa08}. Generation of strong Langmuir turbulence in the downstream region is not easily understood also from the theoretical point of view \\citep[\\eg{},][]{Treumann92,Cairns11}. Moreover, the type-II radio emission itself is also generally observed from the upstream region of interplanetary shock waves, but not from the downstream region \\citep[\\eg{},][]{Reiner98,Bale99}.  There are however observational evidences of radio emission coming from the downstream region of interplanetary shocks \\citep[\\eg{},][]{Hoang92,Lengyel-Frey92,Moullard01}. In addition: 1) some properties of shock waves in the corona may vary from those in the interplanetary medium, 2) spacecraft measurements in the interplanetary space are still limited both by single point measurements and sensitivity of instrumentation, 3) theory of collisionless shocks is still under development and we still do not fully understand them. Also, some suggestions that a shock front has a wavy (rippled) shape, allow us to explain the downstream populations of energetic electrons even in the frame of the standard shock acceleration theory \\citep[\\eg{},][]{Vandas00,Lowe00}. Indeed, anisotropic populations of suprathermal electrons were commonly found downstream from those portion of the Earth's bow shock where the shock normal was quasi-perpendicular to the upstream magnetic field, though suprathermal electrons sharply lost their anysotropy and fluxes with increasing penetration into the sheath \\citep[\\eg{},][]{Gosling89}. This may suggest that populations of nonthermal electrons accelerated at the shock wave front could be also found in the downstream region.  It is not necessary for the non-thermal electron beams responsible for the Langmuir turbulence and the HFC radio emission in the downstream region to be accelerated directly at the shock front. Electrons could also be efficiently accelerated somewhere in the space between the shock front and the leading edge (or on it) of the erupting magnetoplasma structure, \\ie{}, in the magnetosheath (Figure~\\ref{fig9}). For example, it is known both from in situ measurements \\citep[\\eg{},][]{Moullard01,Wei03,Gosling07,Wang10,Chian11} and numerical experiments \\citep[\\eg{},][]{Schmidt03,Wang10} that magnetic reconnection can occur at the interface between the leading edge of interplanetary coronal mass ejections (ICMEs) and background solar wind magnetic field \\citep[see also][]{Demoulin08}. Such episodes of magnetic reconnection could supply beams of suprathermal and/or nonthermal energetic electrons to the sheath region \\citep{Wang10,Huang12}, thus possibly creating the necessary conditions for generation of the Langmuir turbulence and radio emission there. However, we realize (and emphasize) that this important issue requires further studies.  If the upstream-downstream scenario really takes place in the studied event then it is possible \\citep[\\eg{},][]{Smerd74,Smerd75,Mann95,Vrsnak02} to estimate the upstream (\\ie{}, background) magnetic field ($B_{u}$) using the density jump found at the shock front $X=n_{d}/n_{u} \\approx n_{HFC}/n_{LFC}=\\left(f_{HFC}/f_{LFC}\\right)^2=\\left(1+\\left\\langle \\Delta f/f\\right\\rangle\\right)^{2} \\approx 1.35$, and compare it with those values which were estimated in Subsection~\\ref{Versus} using the formula of \\citet{Dulk78}. We will suggest here that the shock wave was oblique rather than purely parallel or perpendicular because a slightly oblique (quasi-perpendicular) shock wave seems to be a more favourable accelerator of electrons in the frame of the shock drift acceleration mechanism \\citep[\\eg{},][]{Holman83}. For an oblique MHD shock wave \\citep[\\eg{},][]{Priest82} with an angle $\\psi$ between the upstream magnetic field and the shock normal it is possible to derive analytically a quadratic equation $aK^{2}+bK+c=0$ which relates the unknown upstream Alfv\\'{e}n-Mach number $M_{A}=\\sqrt{K+X}$ with $\\psi$, $X$ and the Mach number $M_{S}$. Coefficients of this quadratic equation are defined as $a = 6X/M^{2}_{S}+2\\left(X-4\\right)\\cos^{2}\\psi$, $b=X\\left(X+5\\right)\\sin^{2}\\psi$, $c=3X^{2}\\left(X-1\\right)\\sin^{2}\\psi$. Here it was suggested that the adiabatic index $\\gamma=5/3$.  Firstly, let's estimate the Mach number as $M_{S} \\approx v_{LFC} / v_{S}$, where $v_{LFC} \\approx 2.2 \\times 10^{8}$~cm~s$^{-1}$ is the velocity of the LFC source (see Table~\\ref{Table1}) and $v_{S}$ is the upstream sound speed. Within the standard range of coronal temperatures $T \\simeq 1-2$~MK we find $M_{S} \\approx 10.2-14.5$. Now we can find the physically meaningful solution of the quadratic equation and thus estimate the Alfv\\'{e}n-Mach number as $M_{A} \\approx 1.06-1.16$ within the entire range of $\\psi \\in \\left(0, \\pi/2 \\right)$. The upstream magnetic field $B_{u}$ can be estimated as $B_{u} \\approx 4.6 \\times 10^{-12} v_{LFC} n_{u}^{1/2} / M_{A}$~G, where $n_{u}$ is the upstream electron plasma density in cm$^{-3}$ which was already estimated in Subsection~\\ref{Versus} using the observed frequency range of the type-II burst as $n_{u} \\approx 5\\times 10^{7} - 1\\times 10^{9}$~cm$^{-3}$. Thus, we find $B_{u} \\approx 6-33$~G and also the plasma parameter $\\beta = 2 \\left(M_{A}/M_{S} \\right)^2 / \\gamma \\approx 6 \\times 10^{-3} - 13 \\times 10^{-3}$. These values of $B_{u}$ are about six times larger than those obtained in Subsection~\\ref{Versus}. It is not surprising, since the used empirical formula of \\citet{Dulk78} is a generalization of observational data of many different active regions.  A more important inference from the above estimations is the small value of $M_{A} \\approx 1.06-1.16$. It indicates that the upstream Alfv\\'{e}n speed could be $v_{A} \\approx v_{LFC}/M_{A} \\approx 1900-2080$~km~s$^{-1}$, that is significantly larger than the observed speeds of the eruptive plasmas (see Table~\\ref{Table1}). At first glance this may seemed contradictory to the piston-driven shock wave scenario, which, as it was argued in Subsections~\\ref{PDSW} and \\ref{BSW}, is preferable in this event. However, numerical simulations of \\citet[][]{Pomoell08,Pomoell09} have shown that in the inhomogeneous corona even a sub-Alfv\\'{e}nic plasma ejection can launch a shock wave.  It makes sense also to note that the inferred Alfv\\'{e}n-Mach number is less than the critical Alfv\\'{e}n-Mach number $M_{c}\\approx2.76$ for a resistive shock wave \\citep[\\eg{},][]{Treumann09}. Consequently, the shock wave could be subcritical, regardless of whether it was quasi-parallel or quasi-perpendicular. Electrons were definitely accelerated by the shock wave in the event studied since the type-II burst emission was observed. This gives evidences that subcritical shock waves can accelerate electrons in the corona. This fact could be interesting for theories of charged particle acceleration since it is a supercritical shock wave which is generally believed to accelerate charged particles \\citep[\\eg{},][]{Mann95,Treumann09}.   \\subsubsection{Scenario~2} \\label{S2}  In principle, Scenario~2 could also be implemented in the studied event. One of the possible cases is schematically illustrated in Figure~2(b) of \\citet{Holman83}. The efficiency of the type-II radio emission production in the upstream zone by shock drift accelerated electrons depends critically on the angle ($\\psi$) between the shock normal at a given point and the upstream magnetic field. Radio emission from the upstream region is then expected only if $\\psi$ is restricted to a narrow angular range, within a few degrees of $90^{\\circ}$. In the case of a particular mutual arrangement of the shock wave front and of the upstream magnetic field, two separate regions of electron acceleration and thus of enhanced radio emission can be expected. No apparent contradictions between the idea of \\citet{Holman83} (especially illustrated in their Figure~2(b)) and our observations are found. If, in the event studied here, the shock wave front was curved rather than plane, the apparent location of the high frequency component sources (HFC) below the low frequency component (LFC) ones could be mainly due to the projection effect.  Scenario~2 contrarily to Scenario~1 has an important principal drawback --- it can not easily explain correlated intensity and frequency drift variations of the LFC and HFC observed in many type-II bursts \\citep[\\eg{},][]{Vrsnak01}, as well as in the studied event. It also has difficulties to explain the common range of the relative band-split $\\Delta f / f \\approx 0.1-0.2$ in many type-II bursts since the solar corona is very inhomogeneous \\citep[\\eg{},][]{Cairns11}. In our opinion, these facts make Scenario~2 less favourable than Scenario~1.  \\subsubsection{An alternative scenario} \\label{ALI}  \\citet[][]{Cairns94} reported observations of fine-structured electromagnetic emissions both from the solar wind and from the Earth's foreshock. It was shown that for fundamental emission, the fine structures above the local plasma frequency $f_{p}$ corresponded to bands separated by near half harmonics of the electron cyclotron frequency $f_{ce}$, \\ie, by $n f_{ce} / 2$, where $n$ is a natural number. For harmonic emission the separation was $n f_{ce}$.  In our case the frequency separation for the harmonic emission of the observed type-II burst is $\\Delta f \\approx 0.16 f_{LFC}$, \\ie{}, $\\Delta f \\approx 30-60$~MHz. The magnetic field estimated in the previous sections is within the range of $B_{1} \\approx 6-33$~G (using the upstream-downstream hypothesis) or $B_{2} \\approx 1-6$~G \\citep[using the formula of][]{Dulk78}. Consequently, the electron cyclotron plasma frequency should be $f_{ce1} \\approx 17-92$~MHz in the first case and $f_{ce2} \\approx 3-17$~MHz in the second case. The $f_{ce1}$ is in agreement with the observed separation of the LFC and HFC, but $f_{ce2}$ is not. However, the first case corresponds to the situation when the LFC sources were emitted from the upstream region, whereas the HFC sources were emitted from the downstream region. This is in contradiction with \\citet[][]{Cairns94} findings, since he reported that the fine-structured emissions were radiated from the upstream region of the Earth's bow shock only."
    },
    "1208/1208.3677_arXiv.txt": {
        "abstract": "We  study the instrumental features of photons from the peak observed at $E_\\gamma=130$ GeV in the spectrum of Fermi-LAT data.  We use the {\\sc sPlots} algorithm to reconstruct -- seperately for the photons in the peak and for background photons -- the distributions of incident angles, the recorded time, features of the spacecraft position, the zenith angles,  the conversion type and details of the energy and direction reconstruction. The presence of a striking feature or cluster in such a variable would suggest an instrumental cause for the peak.  In the publically available data, we find several suggestive features which may inform further  studies by instrumental experts, though the size of the signal sample is too small to draw statistically significant conclusions. ",
        "introduction": "While the existence of dark matter is widely accepted, its particle nature remains undiscovered.   Potential avenues for discovery include observation of  production at high energy accelerators,  scattering with heavy nuclei in large low-noise underground volumes, or annihilation.  A clear signal of dark matter annihilation may be carried by gamma rays traveling to Earth from regions in the galaxy of high dark-matter density.  As they do not typically scatter after their production, the photon energy and direction are powerful handles for understanding the mechanism of dark matter annihiliation into standard model particles.  One mechanism is annihilation resulting in quarks, which would hadronize and yield $\\pi^0$ particles which in turn produce photons. The spectrum of such a process would give fairly low energy photons ($E_\\gamma \\le\\approx 50$ GeV) which may be difficult to distinguish from other sources.  A more striking feature may appear from annhilation directly into two-body final states including a photon. Rather than yielding a broad energy spectrum,  this process would produce a photon with a well-defined energy given (for the process $\\chi \\chi \\rightarrow \\gamma Y$) by \\bea E_\\gamma = m_\\chi \\left( 1 - \\frac{M_Y^2}{4 m_\\chi^2} \\right) \\label{eq:eline} \\eea where $M_Y$ is the mass of the second annihilation product, such as a $Z$ boson or a second photon.  For the case where $Y=\\gamma$, the line occurs at the mass of the dark matter particle, $E_\\gamma=m_\\chi$.  This makes a search for peaks in the photon spectrum an important component of the dark matter program using Fermi-LAT data\\cite{Abdo:2010nc,Fermi:2012}.  Recent studies have identified a feature in the gamma ray spectrum near $E_\\gamma=130$ GeV\\cite{Bringmann:2012vr,Weniger:2012tx} with a source close to the galactic center \\cite{Bringmann:2012vr,Weniger:2012tx,Tempel:2012ey,finksu}.  The line feature is not accompanied by a lower-energy continuum emission, as would be expected in many models of dark matter interaction~\\cite{wacker}.   However, the large apparant significance of the feature has generated keen interest in exploring other, more mundane explanations, such as unconsidered features in the non-dark-matter background in the difficult region of the galactic center, or instrumental effects in the Fermi-LAT detector.  In this paper, we present a first study of the instrumental characteristics of photons in the line feature, using the {\\sc sPlots}~\\cite{splots} algorithm to disentangle the two populations (background and peak). This allows us to reconstruct distributions in variables which may reveal instrumental issues that would not otherwise be apparant. ",
        "conclusions": "We have performed an initial study of the instrumental characteristics of events from the feature at $E_\\gamma=130$ GeV observed in the Fermi-LAT data.  In the instrumental variables available in the public data distribution, we find no conclusive difference in characteristics between peak photons and background photons, see Table~\\ref{tab:stat}.  There are several suggestive discrepancies, near cos$(\\theta)=0.7$ or  McIlwain $B$ parameter of 1.65 Guass which deserve further study by instrumental experts.  There are several additional instrumental variables which should be examined, such as the incident position on the face of the LAT, but are not available in the public data.  If a striking feature had appeared -- such as a clustering of the peak photons at a given time or near a specific angle of incidence -- it would have pointed to an instrumental issue.  The statistics of the sample are too poor to draw strong conclusions, but the lack of a very clear features makes an instrumental explanation somewhat less likely."
    },
    "1208/1208.1996_arXiv.txt": {
        "abstract": "We search for the presence of double gamma-ray line from unassociated Fermi-LAT sources including detailed Monte Carlo simulations to study its global statistical significance. Applying the Su \\& Finkbeiner  selection criteria for high-energy photons we obtain a similar excess over the power-law background from 12 unassociated sources. However, the Fermi-LAT energy resolution and the present low statistics does not allow to distinguish a double peak from a single one with any meaningful statistical significance. We study the statistical significance of the fit to data with Monte Carlo simulations and show that the fit agrees almost perfectly with the expectations from random scan over the sky. We conclude that the claimed high-energy gamma-ray excess over the power-law background from unassociated sources is nothing but an artifact of the applied selection criteria and no preference to any excess can be claimed with the present statistics. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.4866_arXiv.txt": {
        "abstract": "A cryogenic achromatic half-wave plate (HWP) for submillimetre astronomical polarimetry has been designed, manufactured, tested, and deployed in the Balloon-borne Large-Aperture Submillimeter Telescope for Polarimetry (BLASTPol). The design is based on the five-slab Pancharatnam recipe and it works in the wavelength range 200--600\\,$\\mu$m, making it the broadest-band HWP built to date at (sub)millimetre wavelengths. The frequency behaviour of the HWP has been fully characterised at room and cryogenic temperatures with incoherent radiation from a polarising Fourier transform spectrometer. We develop a novel empirical model, complementary to the physical and analytical ones available in the literature, that allows us to recover the HWP Mueller matrix and phase shift as a function of frequency and extrapolated to 4\\,K. We show that most of the HWP non-idealities can be modelled by quantifying one wavelength-dependent parameter, the position of the HWP equivalent axes, which is then readily implemented in a map-making algorithm. We derive this parameter for a range of spectral signatures of input astronomical sources relevant to BLASTPol, and provide a benchmark example of how our method can yield improved accuracy on measurements of the polarisation angle on the sky at submillimetre wavelengths. ",
        "introduction": "Galactic magnetic fields are believed to play a crucial role in the evolution of star-forming molecular clouds, perhaps controlling the rate at which stars are born and even determining their mass \\citep[][]{Crutcher2004b,McKee2007}. % However, magnetic fields are very difficult to probe on the spatial scales relevant to the star-forming processes, especially within obscured molecular clouds \\citep[e.g.,][]{Crutcher2004a,Whittet2008}, hence their influence on star formation has not yet been clearly established observationally.  Zeeman splitting of molecular lines, which allows a direct measurement of the strength of the line-of-sight component of the local magnetic field, has been carried out successfully for a number of molecular cloud cores \\citep[e.g.,][]{Crutcher1999a}, though the technique is difficult and often limited to very bright regions \\citep{Crutcher2012}. A promising alternative method is to observe clouds with a far-infrared/submillimetre (FIR/submm) polarimeter \\citep[e.g.,][]{Hildebrand1984,Hildebrand2000,WardThompson2000}. By tracing the linearly polarised thermal emission from aspherical dust grains aligned with respect to the local magnetic fields, we can estimate the direction of the plane-of-the-sky component of the field within the cloud \\citep[][]{Davis1951,Dolginov1976,Lazarian2007}, and its strength via the Chandrasekhar \\& Fermi \\citeyearpar[CF;][]{Chandrasekhar1953} technique, provided that ancillary measurements of the turbulent motion velocity are available. The observed morphology in submm polarisation maps can also be used, in synergy with magnetohydrodynamic simulations, to study the imprint of turbulence and magnetization on the formation of structure in the cloud \\citep[][]{Houde2009,Soler2013}.     Ground-based observations with the SCUBA polarimeter \\citep{Murray1997,Greaves2003} and the Submillimeter Polarimeter for Antarctic Remote Observations \\citep[SPARO;][]{Novak2003} show that the submm emission from prestellar cores and giant molecular clouds (GMCs) is indeed polarised to a few percent \\citep{WardThompson2000,Li2006}. {\\sl Planck} \\citep{Planck2011} will provide coarse resolution ($\\sim$5\\arcmin) submm polarimetry maps of the entire Galaxy. The Atacama Large Millimeter/submillimeter Array \\citep[ALMA;][]{Wootten2009} will provide sub-arcsecond millimetre (mm) and submm polarimetry, capable of resolving fields within cores and circumstellar disks, but will not be sensitive to cloud-scale fields.  The Balloon-borne Large-Aperture Submillimeter Telescope for Polarimetry \\citep[BLASTPol;][]{Marsden2008,Fissel2010,Pascale2012}, with its arcminute resolution, is the first submm polarimeter to map the large-scale magnetic fields within molecular clouds with unique combined sensitivity and mapping speed, and sufficient angular resolution to observe into the dense cores. BLASTPol will be able to trace magnetic structures in the cold interstellar medium from scales of 0.05\\,pc out to 5\\,pc, thus providing a much needed bridge between the large-area but coarse-resolution polarimetry provided by {\\sl Planck} and the high-resolution but limited field-of-view maps of ALMA.   BLASTPol successfully completed two science flights over Antarctica during the austral summers of 2010 and 2012, mapping the polarised dust emission at 250, 350, and 500\\,$\\mu$m over a wide range of column densities corresponding to $A_V$\\,$\\gtrsim$\\,4\\,mag, yielding hundreds to thousands of independent polarisation pseudo-vectors per cloud, for a dozen between GMCs and dark clouds. The first scientific results from the 2010 campaign are soon to be released \\citep[][and Poidevin et al. in preparation]{Matthews2013}, while the 2012 data are still under analysis (an overview of the 2012 observations can be found in Angil\\`{e} et al. in preparation).   The BLASTPol linear polarisation modulation scheme comprises a stepped cryogenic achromatic half-wave plate (HWP) and photolithographed polarising grids placed in front of the detector arrays, acting as analysers. The grids are patterned to alternate the polarisation angle sampled by 90$^{\\circ}$ from bolometer-to-bolometer along the scan direction. BLASTPol scans so that a source on the sky passes along a row of detectors, and thus the time required to measure one Stokes parameter (either $Q$ or $U$) is just equal to the separation between bolometers divided by the scan speed. During operations, we carry out spatial scans at four HWP rotation angles spanning 90$^{\\circ}$ (22.5$^{\\circ}$ steps), allowing us to measure the other Stokes parameter through polarisation rotation.  The use of a continuously rotating or stepped HWP as a polarisation modulator is a widespread technique at (sub)mm wavelengths \\citep[e.g.,][]{Renbarger2004,Hanany2005,Pisano2006,Savini2006,Savini2009,Johnson2007,Li2008a,Matsumura2009}. A thorough account of the HWP non-idealities and its inherent polarisation systematics, especially for very achromatic designs, has become necessary as the accuracy and sensitivity of (sub)mm instruments have soared in recent years.  The literature offers numerous efforts to address, through simulations, the impact of the inevitable instrumental systematic errors due to the polarisation modulation strategy in the unbiased recovery of the Stokes parameters $Q,U$ on the sky, especially for cosmic microwave background (CMB) polarisation experiments \\citep[e.g.,][]{Odea2007,Odea2011,Brown2009}. In addition, physical and analytical models have been developed to retrieve the frequency-dependent modulation function of achromatic HWPs and estimate the corrections due to non-flat source spectral indices \\citep{Savini2006,Savini2009,Matsumura2009}.  Nevertheless, little work has been published on incorporating the {\\it measured} HWP non-idealities in a data-analysis pipeline and ultimately in a map-making algorithm. \\citet{Bryan2010} derive an analytic model that parametrises the frequency-dependent non-idealities of a monochromatic HWP and present a map-making algorithm that accounts for these. \\citet{Bao2012} carefully simulate the impact of the spectral dependence of the polarisation modulation induced by an achromatic HWP on measurements of the CMB polarisation in the presence of astrophysical foregrounds, such as Galactic dust. However, both these works assume the nominal design values for the build parameters of the HWP plus anti-reflection coating (ARC) assembly.  While this assumption is a reasonable one when no spectral measurements of the HWP as-built are available, several studies clearly show that the complex multi-slab crystal HWP and its typically multi-layer ARC are practically impossible to manufacture {\\it exactly} to the desired specifications. In particular, \\citet{Savini2006,Savini2009} and \\citet{Pisano2006} caution against the finite precision to which the multiple crystal substrates composing an achromatic HWP can be aligned relative to each other in the Pancharatnam \\citeyearpar{Pancharatnam1955} scheme. In addition, \\citet{Zhang2009} show how some of the design parameters in the ARC can slightly change during the bonding of layers, achieved via a hot-pressing technique \\citep{Ade2006}. We will briefly cover these points and discuss the repercussions on the HWP performance.  This work describes a novel empirical method that allows the reconstruction of the Mueller matrix\\footnote{We adopt the Stokes \\citeyearpar{Stokes1852} formalism to represent the time-averaged polarisation state of electromagnetic radiation; for a review of polarisation basics we refer the reader to \\citet{Collett1993}.} of a generic HWP as a function of frequency through spectral transmission measurements of the HWP rotated by different angles with respect to the input polarised light. Not only does this method give complete and quantitative information on the {\\it measured} spectral performance of the HWP, but it also provides a direct avenue to accounting for the non-idealities of the HWP {\\it as-built} in a map-making algorithm. This empirical approach is applied to the BLASTPol HWP and will help improve the accuracy on astronomical measurements of polarisation angles at submm wavelengths.   The layout of this paper is as follows. In Section~\\ref{sec:intro_HWP}, we give an overview of the manufacturing process for BLASTPol's five-slab sapphire HWP. Section~\\ref{sec:HWP_spectra} describes the spectral measurements, while Section~\\ref{sec:empirical_model} presents the empirical model as well as the main results of the paper. Finally, in Section~\\ref{sec:map_maker}, we describe the algorithm for the naive-binning map-making technique implemented by BLASTPol, which naturally accounts for the {\\it measured} HWP non-idealities. Section~\\ref{sec:concl_HWP} contains our conclusions. ",
        "conclusions": "\\label{sec:concl_HWP}  The goal of the first part of this work was to identify and measure the parameters that fully characterise the spectral performance of the linear polarisation modulator integrated in the BLASTPol instrument, a cryogenic achromatic HWP.  We have described in detail the design and manufacturing process of a five-slab sapphire HWP, which is, to our knowledge, the most achromatic built to date at (sub)mm wavelengths. In the same context, we have provided a useful collection of spectral data from the literature for the sapphire absorption coefficient at cryogenic temperatures.  Using a polarising FTS, we have fully characterised the spectral response of the anti-reflection coated BLASTPol HWP at room temperature and at 120\\,K; we have acquired data cubes by measuring spectra while rotating the HWP to produce the polarisation modulation.  The cold dataset contains measurements in both co-pol and cross-pol configurations; we have used these two data cubes to estimate 9 out of 16 elements of the HWP Mueller matrix as a function of frequency. We have developed an ad-hoc Monte Carlo algorithm that returns for every frequency the best estimate of each matrix element and the associated error, which is a combination of the uncertainty on the measured spectra and a random jitter on the rotation angle.  We have measured how the position of the equivalent axes of the HWP, $\\beta_{\\rm ea}$, changes as a function of frequency, an effect that is inherent to any achromatic design. Once this dependence is accounted for in the Monte Carlo, and a correction is implemented for the residual absorption from sapphire, the Mueller matrix of the HWP approaches that of an ideal HWP, at all wavelengths of interest. In particular, the (band-averaged) off-diagonal elements are always consistent with zero within 2\\,$\\sigma$ and the modulus of the three diagonal coefficients is always $>$\\,0.8. Therefore, we have introduced in the BLASTPol map-making algorithm the band-integrated values of $\\beta_{\\rm ea}$ as an additional parameter in the evaluation of the polarisation angle. To first order, this approach allows us to account for most of the non-idealities in the HWP.  We have investigated the impact of input sources with different spectral signatures on $\\beta_{\\rm ea}$ and on the HWP Mueller matrix coefficients. We find that the HWP transmission and modulation efficiency are very weakly dependent on the spectral index of the input source, whereas the position of the equivalent axes of the sapphire plate stack is more significantly affected. This latter dependence, if neglected, may lead to an arbitrary rotation of the retrieved polarisation angle on the sky of magnitude $2\\,\\overline{\\beta}_{\\rm ea} =10$--15$^{\\circ}$ (3--5$^{\\circ}$) at 250 (500)\\,$\\mu$m. The 350\\,$\\mu$m band, however, is minimally perturbed by this effect.  In principle, the measured Mueller matrix can be used to generate a synthetic time-ordered template of the polarisation modulation produced by the HWP as if it were continuously rotated at a mechanical frequency $f=\\omega\\,t$. Continuous rotation of the HWP allows the rejection of all the noise components modulated at harmonics different than 4\\,$f$ (synchronous demodulation) and is typically employed by experiments optimised to measure the polarisation of the CMB \\citep[e.g.,][]{Johnson2007,Reichborn2010}. In such experiments, the HWP modulation curve leaves a definite synchronous imprint on the time-ordered bolometer data streams, hence it is of utter importance to characterise the template and remove it from the raw data. However, a time-ordered HWP template would be of no use to a step-and-integrate experiment such as BLASTPol, whose timelines are not dominated by the HWP synchronous signal.  We have measured the phase shift of the HWP across the wavelength range of interest to be within 5$^{\\circ}$ of the ideal 180$^{\\circ}$ for the central BLASTPol band, and within 15$^{\\circ}$ for the side bands. This is due to a combination of alignment errors of the sapphire substrates, which are hard to avoid in the manufacture of a five-slab stack, and their lower than ideal thickness. However, the modulation efficiency of the HWP is only mildly affected by this departure from ideality, being above 98\\% in all three BLASTPol bands. Moreover, departures of similar amplitude are not uncommon for HWPs at (sub)mm wavelengths.  The goal of the second part of this work was to include the measured non-idealities of the HWP as-built in a map-making algorithm. We have focused on the implementation of a naive binning technique for the case of BLASTPol, under the assumption of white and uncorrelated noise. As a proof of concept, we have presented a preliminary polarisation map for one of the scientific targets observed by BLASTPol during its first Antarctic flight, completed in January 2011. The inferred direction for the local magnetic field in the Carina Nebula star-forming region is in excellent agreement with the results obtained by \\citet{Li2006} with the SPARO instrument. The empirical approach presented in this paper will help improve the accuracy on astronomical measurements of the polarisation angle on the sky at submm wavelengths."
    },
    "1208/1208.6114_arXiv.txt": {
        "abstract": "Gas within molecular clouds ({\\small MC}s) is turbulent and unevenly distributed.  Interstellar shocks such as those driven by strong fluxes of ionising radiation ({\\small IR}) profoundly affect {\\small MC}s.  While small dense {\\small MC}s exposed to a strong flux of {\\small IR} have been shown to implode due to radiation-driven shocks, a phenomenon called \\emph{radiation driven implosion}, larger {\\small MC}s, however, are likely to survive this flux which in fact, may produce new star-forming sites within these clouds. Here we examine this hypothesis using the Smoothed Particle Hydrodynamics ({\\small SPH}) algorithm coupled with a ray-tracing scheme that calculates the position of the ionisation-front at each timestep. We present results from simulations performed for three choices of {\\small IR}-flux spanning the range of fluxes emitted by a typical {\\small B}-type star to a cluster of {\\small OB}-type stars. The extent of photo-ablation, of course, depends on the strength of the incident flux and a strong flux of {\\small IR} severely ablates a {\\small MC}. Consequently, the first star-formation sites appear in the dense shocked layer along the edges of the irradiated cloud. Radiation-induced turbulence readily generates dense filamentary structure within the photo-ablated cloud although several new star-forming sites also appear in some of the densest regions at the junctions of these filaments. Prevalent physical conditions within a {\\small MC} play a crucial role in determining the mode, i.e., filamentary as compared to isolated pockets, of star-formation, the timescale on which stars form and the distribution of stellar masses. The probability density functions ({\\small PDF}s) derived for irradiated clouds in this study are intriguing due to their resemblance with those presented in a recent census of irradiated {\\small MC}s. Furthermore, irrespective of the nature of turbulence, the protostellar mass-functions({\\small MF}s) derived in this study follow a power-law distribution. When turbulence within the cloud is driven by a relatively strong flux of {\\small IR} such as that emitted by a massive {\\small O}-type star or a cluster of such stars, the {\\small MF} approaches the canonical form due to Salpeter, and even turns-over for protostellar masses smaller than $\\sim$0.2 M$_{\\odot}$. ",
        "introduction": "Stars condense out of cold dense clumps of molecular gas called prestellar cores which appear to be located in much larger clouds, often filamentary in shape. While the formation of stars themselves is relatively well understood, a number of questions such as the distribution of stellar masses or even the formation of prestellar cores remain amongst the unsolved questions in contemporary astronomy. It was commonly believed that stars belonged to larger clusters populated with several tens of other members (e.g. Palla 2005), but the belief is now contested by more recent observations of dozens of clusters (e.g. Guttermuth \\emph{et al.} 2009).  A better understanding of the processes leading to the assembly of potential star-forming gas probably holds the key towards unravelling issues fundamental to the theory of star-formation. Primarily, they are related to the mode of star-formation, in other words clustered as against isolated or a mixture of the two. The origin of the stellar {\\small IMF}, its universal nature and the possible relation with the distribution of core masses is another issue that has baffled researchers. Also a related question is that about the time-scale on which a {\\small MC} is likely to form stars. For a more detailed discussion of these points the interested reader is referred to Hartmann \\emph{et al.} (2011) and Bonnell \\emph{et al.} (2011).  Furthermore, observations in a number of wave-bands, and in particular, the molecular lines at sub mm-wavelengths have now demonstrated the turbulent interiors and the non-uniform distribution of gas within molecular clouds ({\\small MC}s). In fact, a significant fraction of molecular gas appears to reside in dense filamentary regions (e.g., Elmegreen 1997, Nutter \\emph{et al.} 2008, Andr{\\'e} \\emph{et al.} 2010). However, numerical simulations over the last decade have demonstrated with considerable success the tendency of turbulent gas to generate filamentary structure on a relatively short timescale, usually smaller than a sound-crossing time (e.g., Klessen \\emph{et al.} 2000, Jappsen \\emph{et al.} 2005 amongst a number of other works). Though turbulence provides support against self-gravity on a global scale, locally, on scales comparable to the length of driven modes, it assembles gas that can eventually become self-gravitating which probably explains star-formation within larger filamentary structures. Young protostars compete with each-other to acquire mass from their natal pool, a scenario better-known as competitive accretion (Bonnell \\emph{et al.} 2001). Evidently the formation of these dense filaments and their possible fragmentation depends on the prevalent physical conditions within the parent {\\small MC}.  Numerical simulations by Ballesteros-Paredes \\emph{et al.} (1999), Heitsch \\emph{et al.} (2008), and Anathpindika (2009 a,b) amongst those by several other authors, have shown that dense filaments of gas form rapidly via fragmentation of larger pressure-confined gas bodies that may form due to collisions between turbulent flows. The fragmentation itself  is a result of a process called gravoturbulent fragmentation, an interplay between the gravitational instability and hydrodynamic instabilities. These simulations have shown that the star-formation process, beginning with the assembly of gas in dense pockets, once triggered, proceeds rapidly due to the non-linear growth of instabilities. However, the fact that a MC probably spends considerable time in a relatively quiescent state before star-formation is triggered, according to some authors supports the hypothesis of slow star-formation (e.g., Krumholz \\& Tan 2007).  {\\small MC}s in astrophysical environments though, are also exposed to ionising radiation ({\\small IR}), which is likely to profoundly affect their evolution and possibly trigger new episodes of star-formation. Large nebulae such as the Rossette nebula (e.g. White \\emph{et al.} 1997; Poulton \\emph{et al.} 2008), the Eagle nebula with its famous pillars of creation and trunks (e.g. Sugitani \\emph{et al.} 2002), the {\\small M17} nebula (e.g. Hoffmeister \\emph{et al.} 2008), the star-forming cloud {\\small NGC7538} (e.g. Ojha \\emph{et al.} 2004), the nebula {\\small SFO38} in {\\small IC1396} (Choudhary \\emph{et al.} 2010), and cometary globules such as {\\small IC1848} (e.g. Lefloch \\emph{et al.} 1997), and {\\small CG12} (Haikala \\& Reipurth 2010), are just a few examples of star-formation triggered by the {\\small IR}. The Rosette {\\small MC}, for instance, is irradiated by the nearby star-cluster {\\small NGC2244} and  the radiation induced turbulence appears to have generated a network of filaments within the main cloud. More recent observations of the Rosette {\\small MC} using the multi-waveband infrared cameras, PACS and SPIRE, on-board the Herschel space-observatory have revealed young star-forming regions in the junctions of these filaments (Schneider \\emph{et al.} 2012).  Secondary star-formation, i.e., cases of new episodes of star-formation triggered by the flux of {\\small IR} emitted by an earlier generation of stars have also been reported in for e.g., the Orion {\\small MC} (Wilson \\emph{et al.} 2005). Another prognosis, observed sometimes in regions of triggered star-formation, is the increasing age of stellar population along a certain direction, also called the sequential mode of star-formation (e.g., Maaskant \\emph{et al.} 2011). Hot radiation from young stars is therefore believed to play a crucial role in controlling the rate of galactic star-formation, and furthermore, it is also an important source of turbulence in galactic disks (e.g. Andrews \\& Thompson 2011; Krumholz \\& Matzner 2009).  Over the last three decades several authors have studied analytically the likely fate of {\\small MC}s irradiated by a strong flux of {\\small IR}. Dyson (1973), for example, suggested flattening of an irradiated cloud, whereas according to Bertoldi (1989) and Bertoldi \\& McKee (1990), an irradiated cloud was more likely to implode under the influence of a radiation driven shock-wave within the cloud, a phenomenon known as \\emph{radiation driven implosion}. The cloud so accelerated was shown to acquire a cometary structure with a characteristic velocity of the order of a few km/s. Stars could form within this comet-shaped globule ({\\small CG}) . Indeed, a number of such globules have been observationally reported as cited above. The formation of {\\small CG}s due to the {\\small RDI} has been demonstrated numerically for small dense clouds by for instance, Lefloch \\& Lazareff (1994), and more recently by Bisbas \\emph{et al.} (2011).  Dale \\emph{et al.} (2007), on the other hand, showed that the exposure of a massive self-gravitating, turbulent {\\small MC} to a flux of strong {\\small IR} significantly reduced the gas-to-star conversion efficiency. While these authors considered only one choice of the strength of {\\small IR}, in the present study we will consider three choices ranging from a weak to a fairly strong flux, similar to that emitted by a typical young star-cluster. The different strengths of flux are acquired by varying the temperature of the source of {\\small IR} and the rate of photon-emission, as recorded in Table 1. We will then compare the evolution of an irradiated, turbulent cloud with that of a similar cloud allowed to evolve only under self-gravity (and no {\\small IR}). Of particular interest is the spatial distribution of dense pockets of gas within the {\\small MC} in each realisation. The plan of the paper is as follows. We shall begin by briefly discussing the expected fate of irradiated MCs for different choices of gas density.  The numerical scheme adopted for this work will be described briefly  in \\S 3, and the simulations will be discussed in \\S 4 before results are formally presented in \\S 5. We will conclude in \\S 6.  \\section[]{Evolution of an irradiated cloud} The physical situation is illustrated in Fig. 1, where the projection (on the plane of the sky) of the cloud facing the external source of radiation, S, and located a distance $r_{s}$ from S have been shown. The flux of {\\small IR}, $d\\Sigma$, received by a small area element on the surface of the cloud as that shown by the small shaded portion in Fig. 1, located at a distance $r$ from the source, S, is \\begin{equation} d\\Sigma = \\frac{\\mathcal{N}_{LyC}}{4\\pi r_{s}^{2}}\\Big(\\frac{R_{cld}}{r}\\Big)^{2} \\sim \\frac{\\mathcal{N}_{LyC}}{4\\pi r_{s}^{2}}\\Big(\\frac{R_{cld}}{r_{s}}\\Big)^{2} \\ \\ \\mathrm{cm^{-2}\\ s^{-1}}, \\end{equation} where higher order terms involving $(R_{cld}/r_{s})$ have been neglected as $r_{s}\\gg R_{cld}$ in general. The force, $F_{r}$, exerted by this flux of {\\small IR} on the cloud is then \\begin{displaymath} F_{r} = \\pi R_{cld}^{2}\\int_{\\alpha=0}^{\\alpha=\\pi/2}\\frac{c^{2}d\\Sigma}{\\xi^{(2)}n_{i}r_{s}}\\sin(\\alpha)d\\alpha\\ \\mathrm{cm\\ s^{-2}}, \\end{displaymath} where $c$ is the speed of light, $n_{i}$, the average number density of ionised gas and $\\xi^{(2)}$, the second recombination coefficient defined as \\begin{equation} \\xi^{(2)} = \\frac{2.06\\times 10^{-11}Z^{2}}{[T_{ion}/\\mathrm{K}]^{1/2}}\\phi_{2}(\\beta)\\ \\mathrm{cm}^{3} \\mathrm{s}^{-1}. \\end{equation} In the equation above, $\\beta = \\frac{1.58\\times 10^{5}Z^{2}}{[T_{ion}/\\mathrm{K}]}$, $T_{ion}$ is the equilibrium temperature of the ionised gas, and $\\phi_{2}(\\beta)$ is the second recombination coefficient (Spitzer 1978). ($Z$= 1, for a fully ionised gas).  Using Eqn. (1), the expression for $F_{r}$ above becomes \\begin{equation} F_{r} = \\frac{c^{2}\\mathcal{N}_{LyC}}{4n_{i}\\xi^{(2)}r_{s}}\\Big(\\frac{R_{cld}}{r_{s}}\\Big)^{4}\\ \\mathrm{cm\\ s^{-2}}, \\end{equation} Note that the denominator in this expression for radiation-force has units of velocity that may be interpreted as the rate at which the ionisation-front advances. As expected, the force becomes vanishingly small at large distances, $r_{S}$. Equation (3) for the force exerted by the flux of incident radiation permits us to define the momentum, $p_{r}$, delivered  by this flux to the gas within the cloud; \\begin{displaymath} F_{r} = \\frac{dp_{r}}{dt}. \\end{displaymath} For a molecular cloud composed of the usual cosmic mixture the average mass of gas particles within it, $\\bar{m}$ =4$\\times 10^{-24}$ g. If the irradiated cloud loses gas at roughly constant velocity, $v$, then the rate of mass-loss, $\\dot{\\mathrm{M}}_{loss}(t)$, for this cloud can be obtained by applying the conservation of momentum to the mass lost, M$_{loss}(t)$, so that \\begin{equation} \\dot{\\mathrm{M}}_{loss}(t)\\equiv\\frac{d\\mathrm{M}_{loss}(t)}{dt} \\sim \\frac{1}{v}\\frac{dp_{r}}{dt}\\sim 0.25\\frac{\\bar{m}c^2\\mathcal{N}_{LyC}}{r_{s}n_{i}\\xi^{(2)}v}\\Big(\\frac{R_{cld}}{r_{s}}\\Big)^{4}. \\end{equation}  The expression of the mass-loss rate here assumes that the density, $n$, within the cloud is uniform and that there is no secondary loss of the {\\small IR}-flux due to absorption by dust within the cloud. In a real cloud though, a fraction of the incident flux would be lost towards heating the dust within the cloud. The heated dust will of course, re-emit in infrared wavebands. The assumption of a pristine cloud devoid of any dust, however, is unlikely to alter significantly the order-of-magnitude estimate provided by Eqn. (4). The timescale, $\\tau$, over which an irradiated cloud is likely to survive is simply, \\begin{equation} \\tau = \\frac{M_{cld}}{\\dot{\\mathrm{M}}_{loss}(t)}. \\end{equation}  Using Eqn. (4) we calculate the mass lost by irradiated clouds of different average densities, $n$, over their respective free-fall times. The results of a demonstrative calculation performed for a photon-ionised gas of average density, $n_{i}\\sim$ 10 cm$^{-3}$, and maintained at an equilibrium temperature, $T_{ion}$= $10^{4}$ K, have been plotted in Fig. 2. The test cloud for each choice of flux-strength was kept at a distance $r_{s}$ from the source such that $R_{cld}/r_{s}$= 0.01. Starting from the lowest characteristic upward, the average gas density within a cloud, $n$, increases by an order of magnitude. The three green vertical lines for the respective choice of $\\mathcal{N}_{LyC}$ divides the characteristics into survival-regions, i.e., the range of {\\small IR}-fluxes that a cloud can possibly endure. The minimum mass of a cloud that can survive without being photoevaporated by the incident radiation-flux is defined by the intersection of these vertical lines with the characteristics on the plot.  Small rarefied clouds sit leftward of the vertical line for $\\mathcal{N}_{LyC}\\ =10^{48}$ s$^{-1}$ and therefore, are likely to be photo-evaporated even by a relatively weak flux of {\\small IR}. Denser and more massive clouds, to the right of this line, will be steadily photo-evaporated by a stronger flux. Photoablation is therefore the likely fate of progressively larger clouds exposed to a flux of ionising radiation. Some of the more massive clouds occupy the top right-hand corner of the plot to the right of the vertical line for $\\mathcal{N}_{LyC}$ = 10$^{51}$ s$^{-1}$, and evidently need stronger fluxes before they can start losing their mass.  The test cloud used in the simulations below ($n\\sim 10^{4}$ cm$^{-3}$,$R_{cld}$ = 1 pc; see \\S 3.4 below), lies on the characteristic just below the uppermost in the region between the fluxes corresponding to 10$^{49}$ s$^{-1}$ and 10$^{51}$ s$^{-1}$, and therefore, is expected to steadily photo-ablate. In fact, the mass-loss rate for this test cloud according to Eqn.(4) above is a few times $10^{-4}$ M$_{\\odot}$ yr$^{-1}$ for $\\mathcal{N}_{LyC}\\sim\\ 5\\times10^{49}$ s$^{-1}$,  and $R_{cld}/r_{s}\\sim$ 0.08. The mass-loss rate for this cloud suggests, it is likely to survive at least an order of magnitude longer than its free-fall time.  \\begin{figure} \\vspace{20pt} \\includegraphics[width=7cm, angle=0]{IONSNCLD.eps} \\caption{A  projection showing the location of the molecular cloud ({\\small MC}), of radius R$_{cld}$, relative to the source of ionisation, $S$. The amount of flux received by the shaded area element on the surface of the cloud is given by Eqn. (1) in the main text. } \\end{figure}    \\begin{figure} \\vspace{20pt} \\includegraphics[width=7cm, angle=270]{IRcldclassf.eps} \\caption{A simple classification scheme for {\\small MC}s that enables us to predict its evolution when exposed to a flux of {\\small IR}; the calculation assumes $n_{i}$ = 10 cm$^{-3}$, $T_{ion}$ = 10$^{4}$ K and $r_{s}/R_{cld}$=12. The three choices of $\\mathcal{N}_{LyC}$ used in the present work are shown by green vertical lines. The intersection of these lines with the cloud mass-flux characteristics defines the minimum mass required for a cloud of given density to survive the incident flux of {\\small IR}. \\emph{See text for description.}} \\end{figure}   \\section[]{Numerical Method} \\subsection{Smoothed Particle Hydrodynamics ({\\small SPH})} The {\\small SPH} algorithm was first introduced by Gingold \\& Monaghan (1977), and Lucy (1977) to handle complex problems in astrophysical context and other areas of computational fluid dynamics (e.g. Monaghan 2005). We use our numerical hydrodynamics code, {\\small SEREN}, an open-{\\small MP} parallised algorithm that has been extensively tested for numerous applications of astrophysical interest (Hubber \\emph{et al.} 2011). The fundamental quantity in an {\\small SPH} code, the density of particles, is calculated using the \\emph{gather} technique. The search for nearest neighbours of an {\\small SPH} particle, and the calculation of the net force (gravity+hydrodynamic) on an {\\small SPH} particle is done by distributing particles on the Barnes-Hut tree, with a cell-opening angle of $\\sim$0.2; quadrupole moments of distant cells in the tree structure are also included in the calculation of forces. We have used the standard prescription for {\\small SPH} artificial viscosity with the viscous parameters having values (0.1,0.2).  The effects of ionising radiation are included through a ray-casting algorithm devised by Bisbas \\emph{et al.}(2009), this scheme is an inherent  part of {\\small SEREN}. The basic features of this algorithm are : it - (a) determines the position of the ionisation front ({\\small IF}) at a given instant of time, and (b) assigns appropriate temperatures to the {\\small SPH} particles. Of the two, (a) is achieved via the {\\small HEALP}ix algorithm (G{\\'o}rski \\emph{et al.} 2005). Individual rays originating from the source of {\\small IR} that trace the direction of propagation of the radiation are distributed uniformly on a sphere by {\\small HEALP}ix. Rays are distributed at different levels, $l$, between 0 and 7 so that there are $12\\times 4^{l}$ rays on any level. The angular resolution of the {\\small IF} is determined by the value of $l$. The maximum intensity, $I_{max}$, at the {\\small IF} from Eqn. (6) is \\begin{equation} I_{max}\\equiv \\frac{m^{2}\\mathcal{N}_{LyC}}{4\\pi\\xi^{(2)}}, \\end{equation} where $m=\\frac{\\bar{m}}{X}$, is the mean mass of a gas molecule and $X$=0.7, and $\\xi^{(2)}$ is the second recombination coefficient defined by Eqn. (2) above.  The integral form of the above expression for intensity at a point $r$ is, \\begin{equation} I(r) = \\int_{r} \\rho^{2}(r\\mathbf{\\hat{e}})r^{2}dr, \\end{equation} where $\\mathbf{\\hat{e}}$ is the unit vector in the direction of $r$. Equation (7) is solved iteratively to determine $r$ such that $I(r)=I_{max}$, while calculating the {\\small SPH} density at each point of the iteration. \\subsection{Thermodynamics} All points satisfying the inequality $I(r) < I_{max}$, lie within the HII region and are therefore assigned a temperature, T$_{ion}$, the equilibrium temperature of the HII region. While the equilibrium temperature probably depends on the coupling between matter and radiation (e.g., Genel \\emph{et al.} 2012), in the present study, T$_{ion}$, is assigned a fiducial value of 8000 K. Neutral gas within the irradiated cloud is treated isothermally and maintained at a temperature, T$_{n}$ = 20 K. The algorithm also includes a temperature smoothing scheme that alleviates the discontinuity arising due to a sharp difference in the temperature of the ionised and neutral gas particles. The smoothing scheme is essentially a Taylor-series approximation to the first order.  \\begin{table} \\centering \\begin{minipage}{80mm} \\caption{Listed below are the choices of ionising flux and the temperature of the source used in various test cases (Spitzer 1978). The temperature of the ionised gas is maintained at, T$_{ion}$ = 8000 K in cases 2, 3 and 4. } \\begin{tabular}{@{}rcc@{}} \\hline \\hline Serial & $\\frac{\\mathcal{N}_{LyC}}{[\\mathrm{s}^{-1}]}$; $\\frac{\\mathrm{T}_{star}}{[K]}$ & $\\Sigma_{\\mathcal{N}_{LyC}}$ \\\\ No.  &   &  [cm$^{-2}$s$^{-1}$]\\\\ \\hline 1  & No ionising radiation & \\\\ \\hline 2   & 4$\\times 10^{48}$;  &  $1.10\\times 10^{7}$  \\\\ &  36,000 K &  \\\\ \\hline 3  & 5$\\times 10^{49}$;   &  $1.37\\times 10^{8}$ \\\\ & 47,000 K & \\\\ \\hline 4  & 7$\\times 10^{51}$;  &  $1.93\\times 10^{10}$  \\\\ &  51,000 K & \\\\ \\hline \\end{tabular} \\end{minipage} $\\mathcal{N}_{LyC}, T_{star}, \\Sigma_{\\mathcal{N}_{LyC}}$, are  respectively the rate of ionising photons, the temperature of the source star and the flux of ionising radiation \\end{table}  \\subsection{Sink particles} An {\\small SPH} particle with density higher than the threshold, $\\rho_{thresh} \\sim 10^{-15}$ g cm$^{-3}$, is replaced by a special type of particle, the sink, following the prescriptions made by Bate \\& Burkert (1997). Apart from the density threshold, a sink particle is also has a radius which in the present simulations is set as 2.5 times the {\\small SPH} smoothing length of a prospective sink particle. {\\small SPH} particles bound gravitationally to the sink are accreted by it. The sink particle in the present numerical exercises represent a protostellar object. The minimum mass of an {\\small SPH} particle, M$_{min} \\ = \\ {\\mathrm N}_{neibs}\\Big(\\frac{\\mathrm{M}_{cld}}{\\mathrm{N}_{tot}}\\Big)$, where $\\mathrm{N}_{neibs}\\ = \\ 50$ and $\\mathrm{N}_{gas}\\ =\\  550,000 $, are respectively the number of nearest neighbours of an {\\small SPH} particle and the number of particles within the cloud. With this choice of {\\small SPH} parameters the ratio of the thermal Jeans mass, $\\mathrm{M}_{J}(\\bar{\\rho}\\sim 10^{-16}\\mathrm{g\\ cm}^{-3},T_{n}=20\\ \\mathrm{K})$ against the minimum mass, M$_{min}$, is 6 which satisfies the Bate-Burkert criterion of resolving the gravitational instability.  \\subsection{Initial conditions} The test cloud used in the present study was modelled as a uniform-density sphere of unit mass and radius and was assembled by randomly positioning particles within it. Since the randomness in positioning particles is inherently associated with Poissonian noise, we first evolved it for a fraction of a sound-crossing time allowing the  numerical viscosity to dissipate the spurious noise. This procedure settled the initial assembly of particles which was then rescaled to the desired dimensions. The mass and radius of the cloud are respectively, $M_{cld} = 400$ M$_{\\odot}$, and radius, $\\mathrm{R}_{cld}$ = 1 pc, where as the gas within it was maintained at a uniform temperature, ${\\mathrm T}_{cld}$ = 20 K. A supersonic, Gaussian velocity field with random amplitudes and a relatively steep power-spectrum, $P(k)\\propto k^{-4}$,  was then superposed on this cloud \\footnote{$k \\& P(k)$, are respectively the wave-vector and the power in the wave-vector space}. This velocity field had an initial Mach number of 10 and was first set up on a grid with 128$^{3}$ cells in the $k$-space. The test cloud was then mapped on the $k$-space and velocity components for individual particles were calculated via interpolation. Then a random phase was added to each velocity component of a particle while transforming it back into the Cartesian-space. The cloud with its turbulent velocity field was then allowed to evolve under self-gravity; this is the first case listed in Table 1. This cloud soon developed filamentary structure as shown in the rendered density plot of Fig. 3 ($t\\sim 0.05$ Myr), which was then used as the initial condition in the remaining 3 cases, listed 2 through to 4 in Table 1. The cloud age was not reset to zero at the time of introducing the source of {\\small IR} in these latter cases.  Each test case discussed here took a little over 5000 CPU hours and was run on the {\\small HYDRA} supercomputing cluster at the Indian Institute of Astrophysics. The cluster is composed of the Intel Xeon-5675 processor. Simulations were allowed to form 150 sink particles before they were terminated.   \\begin{figure} \\vspace{20pt} \\includegraphics[angle=270, width=8cm]{IR_ics.eps} \\caption{A rendered density plot, at $t = 0.2t_{ff}\\sim$ 0.05 Myr, showing the mid-plane of an initially turbulent MC. Injected turbulence readily generates dense filaments within the cloud, as can be seen in this plot. This cloud is then used as the initial condition for the remaining test cases when the source of IR is turned on ($t_{ff}\\sim$ 0.27 Myr).} \\end{figure}  \\section[]{Numerical Simulations} \\subsection{Turbulent cloud without a source of ionising radiation}  \\begin{figure} \\vspace{20pt} \\includegraphics[angle=270, width=9.cm]{GRmontg.eps} \\caption{A sequel of column density plots showing the mid-plane of the turbulent cloud at different epochs; $t$ = 0.1,0.15, 0.2 $\\&$ 0.25 Myr, for images in respective panels from the top to bottom ($t_{ff}\\sim$ 0.27 Myr). Turbulence within the cloud generates filamentary structure that can be seen in the figure on each panel.} \\end{figure}  Collisions between random gas flows within a turbulent cloud dissipate turbulent energy and generate filamentary structure within the cloud as can be seen in the renders plotted in Figs. 3 and 4.  The initial velocity field, with a power-spectrum considerably steeper than the Kolmogorov type, was injected between wavenumbers 1 and 8. As a result turbulent energy cascades down via shocks from larger spatial scales. Respective plots in the panels of Fig. 4 show the rendered density images of the mid-plane of the cloud at successive epochs. After a little over one free-fall time the original cloud collapses to a centrally located dense filament with a few arterial extensions. The advanced stage of the collapsed cloud can be seen more closely in the rendered plot of Fig. 5 where positions of a few protostellar objects have also been marked. As expected, the formation of protostars proceeds along the dense filaments.  \\begin{figure} \\vspace{20pt} \\includegraphics[angle=270, width=8cm]{GRfin.eps} \\caption{A column density plot showing an advanced stage in the evolution of the turbulent cloud in case 1; $t$ = 0.31 Myr$\\sim$ 1.2 $t_{ff}$. The position of the sink particles within the collapsed cloud has been marked with * on this plot. Note that sinks represent protostellar objects.} \\end{figure}  \\subsection{Irradiated turbulent cloud} \\begin{figure} \\vspace{20pt} \\includegraphics[angle=270, width=9.cm]{WIONSNmontg.eps} \\caption{The mid-plane of the cloud in case 2 irradiated by a weak flux of {\\small IR} is shown in this sequel of column density plots. A faint trunk that eventually evaporates due to the flux of {\\small IR} can be seen in the lower left corner of the second panel. While exposure of the cloud to the incident flux of {\\small IR} causes the cloud to loose mass, the radiation-induced shocks produce dense filamentary structure within the cloud; especially see the plot in the lower panel ($t$ = 0.216 Myr $\\sim$0.8 $t_{ff}$).} \\end{figure}  Exposure to a flux of ionising radiation commonly causes the molecular cloud to lose mass via photo-ablation. Unlike small, dense clumps which when exposed to {\\small IR} implode to form comet-shaped globules, the larger clouds such as the one considered in the present study develops sporadic pockets of dense gas, often filamentary and eventually spawn stars. In this next set of 3 simulations we endeavour to study the effect of varying the strength of {\\small IR} on the evolution of a {\\small MC}.  \\begin{figure} \\vspace{20pt} \\includegraphics[angle=270, width=8cm]{WIONSNfin.eps} \\caption{A column density plot showing the picture in the lowest panel of Fig. 6. Dense filaments within the irradiated cloud are readily visible in this plot and a few sink particles (i.e., protostellar objects) have been marked with *, which unsurprisingly, are located along the filaments ($t\\sim$ 0.24 Myr).} \\end{figure}  \\begin{figure} \\vspace{20pt} \\includegraphics[angle=270, width=8cm]{IRmassloss.eps} \\includegraphics[angle=270, width=8cm]{massthresh.eps} \\caption{\\emph{Top-panel} : Plot showing the fractional mass, $f$, of the irradiated cloud retained in each of the three realisations. \\emph{Lower-panel} : The mass of gas within the cloud above the density threshold, $\\rho_{thresh}\\sim 10^{-18}$ g cm$^{-3}$, has been shown at different epochs of time. Relatively stronger shocks, such as the ones in cases 3 and 4, assemble material in the dense phase on a much shorter timescale.} \\end{figure}   \\subsubsection{Weak ionising radiation} The exposure of a {\\small MC} to a relatively weak flux of {\\small IR}, listed as case 2 in Table 1, produces a C-shaped ionisation front ({\\small IF}), on the surface of the cloud exposed to the source of {\\small IR}. The incident flux of {\\small IR} also shocks the pre-existing filaments within the cloud, visible in the plots on the upper two panels of Fig. 6. Interestingly, the shocking of a clump in one of these filaments also generates a small, faint elephant trunk-like extension that can be seen in the lower left-hand corner of the second panel of this figure. The growth of this trunk is probably the result of a backward re-expansion of material from a shocked clump in a filament within the cloud. This possibility has previously been demonstrated in numerical simulations by for instance, Williams \\emph{et al.}(1999), and more recently by Gritschneder \\emph{et al.} (2009,2010). Protruding trunks are commonly found in star-forming clouds and some even harbour a YSO in the head region such as those reported in the famous M16 (Eagle) nebula (e.g., White \\emph{et al.}1999). However, the trunk observed here ablated soon after its appearance. Turbulence injected by the flux of {\\small IR} further assists the formation of filaments within the cloud, the late phase of which, with a network of filaments and star-forming sites within it can be seen in the rendered density plot in Fig. 7. Also marked on this rendered plot are the positions of protostars that appear aligned with dense gas filaments.  The fraction, $f\\equiv \\mathrm{M}_{ret}(t)/\\mathrm{M}_{cld}$, of the irradiated cloud retained has been plotted in the upper panel of Fig. 8. The fraction of the cloud lost is therefore $(1-f(t))$. Exposure to a weak flux of IR causes the cloud in this case to lose $\\sim$30\\% of its original mass and so a large portion of the cloud is still retained. The residual cloud evolves like the turbulent cloud discussed in \\S 4.1 above, although the network of filaments within it, as can be seen in Fig. 7 and the lowest panel of Fig. 6, is significantly more branched in stark contrast to the filament seen in Fig. 5. The incident flux of a weak {\\small IR} appears to drive many more modes. The simulation was terminated after about 0.9 free-fall times ($t$ = 0.24 Myr), by which time 150 protostars had formed within the cloud.   \\subsubsection{A relatively strong ionising radiation}  With a stronger flux of {\\small IR}, the {\\small MC} in this case, listed 3 in Table 1, loses mass at a much higher rate than that observed in the previous realisation. Consequently, after $t=0.8 t_{ff}\\sim 216$ kyrs, a little over 40 \\% of the original cloud is lost via photo-ablation. Filaments in the original cloud, as is evident from the sequel of rendered density plots in Fig. 9, are shocked by the incident flux of {\\small IR}. This flux also produces a thin, dense shocked-shell on the surface of the irradiated cloud. The surface of the {\\small IF} as can be easily seen, is corrugated. A closer examination reveals its wiggly nature which resembles features of the well-known thin-shell instability ({\\small TSI}). The {\\small TSI}, as the name suggests, is triggered in thin shells confined by ram-pressure either on one or both faces. The instability grows due to the transfer of momentum between perturbed regions of the shell as demonstrated by Vishniac (1983), through a detailed perturbation analysis of the problem. The dynamically unstable nature of the irradiated surface has also been discussed by Williams (2002), who demonstrated the appearance of corrugations on the surface of the {\\small IF}. However, it is unclear if these later corrugations are also the manifestations of the {\\small TSI}, though for the moment we do not attempt to distinguish between the two.  We therefore use a previously derived expression for the length of the fastest growing mode, $\\lambda_{fast}$, to verify if the instability is indeed resolved in this simulation. The length of the fastest growing mode in a shocked layer of uniform surface density, $\\sigma_{0}$, calculated by Anathpindika (2010) is \\begin{displaymath} \\lambda_{fast} = \\frac{2\\pi}{k_{fast}}, \\end{displaymath} where \\begin{equation} k_{fast} = \\frac{\\pi G\\sigma_{0}}{\\Big[a_{0}^{2} - \\frac{P_{E}}{\\rho_{s}}\\Big(1-\\frac{R_{s}}{R_{2}}\\Big)^{-1}\\Big]}. \\end{equation} This mode grows on a timescale, $t_{growth}\\sim \\lambda_{fast}/a$; $a=(k_{B}T_{gas}/\\bar{m}_{H})^{1/2}$ is the sound speed  and, $T_{gas}$, the average temperature of the gas in this layer, which is $\\sim 10^{4}$ K. The radius of the irradiated cloud, $R'_{2} = R_{cld} + dR$, where $R_{cld}$ is the radius of the original cloud and $dR$ is the thickness of the shell. In the present case,  $dR/R'_{cld}\\sim 10^{-2}$, and  $P_{E}/\\rho_{s}\\sim a_{HII}^{2}$ is the square of the sound-speed in ionised gas. Plugging in the appropriate values we get, $\\lambda_{fast}\\sim $0.05 pc, and $t_{growth}\\sim $0.22 Myr, which agrees with the timescale on which wiggles first appear on the surface of the shocked-shell.  The ratio of $\\lambda_{fast}$ to the average {\\small SPH} smoothing length, $\\mathcal{X}$, is a good indicator of the spatial resolution, which in the present case is 3. It is therefore likely that the instability has been only barely resolved here. Better resolution of the instability demands at least an order of magnitude increase in $\\mathcal{X}$, which transforms in to about 3 orders of magnitude increase in the number of {\\small SPH} particles. This is computationally demanding. However, since we are not concerned about the detailed modelling of this instability, the present choice of spatial resolution is sufficient. It is therefore clear that radiation-induced shock produces pockets of dense gas along the irradiated surface, and further, a few other isolated pockets of dense gas can also be seen along the shocked filaments in the interior of the cloud. The rendered density plot in the left-hand panel of Fig. 10 shows the irradiated cloud more closely along with the location of sink particles on it. As with the earlier case, this simulation was also terminated at $t\\sim 0.22$ Myr, by which time the cloud had formed 150 sink particles. We note that the sink-formation timescale in cases 2 and 3 is therefore mutually comparable.  \\subsubsection{Strong ionising radiation} In this final test case, listed number 4 in Table 1, the flux is the strongest with source characteristics similar to those of a typical young star-cluster. The irradiated cloud in this case evolves in a manner similar to that observed in the previous case, albeit on a much shorter timescale and has a much higher mass-loss rate, as also reflected by the plots in Fig. 8. Analogous to the previous 2 cases, exposure to a flux of {\\small IR} shocks the surface and the interiors of the cloud. Some of the clumpy features near the surface briefly appear to develop trunk-like features that  are stymied soon due to rapid photo-ablation of the surface. The surface of the irradiated cloud, as in case 3, also shows evidence of the thin-shell instability discussed previously in \\S 4.2.2. Since the broad features of the irradiated cloud in this case match with those in the previous realisation, we only show the final state of the cloud in the right-hand panel of Fig. 10. Unlike the irradiated cloud in the previous case the initial star-formation in this case is confined predominantly to the shocked shell   \\begin{figure*} \\vbox to 70mm{\\vfil \\includegraphics[angle=270, width=16.cm]{IRSTLRINmontg.eps} \\caption{Mid-plane of the irradiated cloud in case 3 has been shown in these rendered column density plots. Exposure to a relatively strong flux of IR ablates the cloud while also shocking the dense filaments within it as can be seen in the plots on each panel; $t\\sim 10^{5}$, 2$\\times 10^{5}$ and 2.1$\\times 10^{5}$ yrs ($\\sim 0.7 t_{ff}$), for plots in the left-, central and the right-hand panels respectively. } \\vfil} \\label{landfig} \\end{figure*}  \\begin{figure*} \\vbox to 90mm{\\vfil \\includegraphics[angle=270, width=8.cm]{STLRIONSNfin.eps} \\includegraphics[angle=270, width=8.cm]{CLSTRIONSNfin.eps} \\caption{\\emph{Left-panel}: A rendered column-density plot showing the mid-plane of the irradiated cloud in case 3 ($t\\sim 0.75 t_{ff}$), after the formation of the first few protostellar objects. The protostars, represented by the sink-particles, have been marked with '$\\ast$'. Observe that the formation of stars commences in the shocked layer on the cloud surface, though another potential site of star-formation appears in the central region of the irradiated cloud. \\emph{Right-panel}: A rendered column density plot showing the mid-plane of the irradiated cloud in case 4. As in case 3, most of the star-formation in this case is also confined to the shocked layer on the irradiated surface of the cloud ($t\\sim 0.1 Myr\\equiv 0.33t_{ff}$). Shocked filaments within the cloud, thus far, do not appear to have acquired sufficient density to spawn stars. } \\vfil} \\label{landfig} \\end{figure*} ",
        "conclusions": "In this article we presented a comparative study of the evolution of turbulent {\\small MC}s. Turbulence, irrespective of whether it is damped or injected continuously, generates filamentary structure within {\\small MC}s. In the first test case of this study, turbulence was allowed to damp without ever being replenished where as in the remaining three test cases, it was injected continuously via radiation-driven shocks. These latter three cases show that irradiating a large {\\small MC} profoundly affects its evolution, and therefore, the star-formation history. The cloud with damped turbulence, however, evolves significantly differently compared to that in which turbulence is continuously injected via an external source. Not only is there a difference in the evolution, but the timescale on which the cloud evolves also varies and in fact, the cloud with stronger turbulence within it evolves much faster. Thus in the present work, the injected turbulence within the cloud is the strongest in case 4  and so, star-formation within the cloud in this case commences on the shortest timescale in comparison to the other 3 cases.  Global gravity soon dominates gas dynamics within the cloud evolving purely under self-gravity so that few modes other than the Jeans unstable mode are driven. Consequently the Jeans mode becomes the principal mode of instability so that gas within the {\\small MC} is primarily channelled into this mode. Furthermore, this tendency is reflected by the {\\small PDF} as well as the {\\small MF} for the {\\small MC} in case 1. Without much surprise, the {\\small PDF} for this {\\small MC} is largely predisposed towards somewhat higher masses, an obvious mis-fit to the canonical {\\small IMF}. Driven turbulence, on the other hand, injects energy in to various unstable modes which aids formation of dense structure. In this study turbulence was injected within {\\small MC}s by an uninterrupted source of {\\small IR}. Simulations were performed by varying the strength of the {\\small IR}-flux, spanning a range of radiation emitted by a typical {\\small B}-type star to that emitted by young {\\small OB}-associations. There is little doubt that progressively stronger radiation drives shocks of ever increasing strength.  A common feature of the irradiated {\\small MC}s observed in each of the three test cases, 2, 3 and 4 of this study, is the appearance of a shocked shell on the irradiated surface facing the source of IR and the shocked internal filamentary structure generated by the initial turbulence. The secondary shocks driven by the flux of {\\small IR} create new sites of star-formation, particularly in the shocked shell and junctions of filaments within the {\\small MC}. Evidently shocks distribute gas over a wider range of density as reflected by the corresponding {\\small PDF}s. The range over which gas-density is distributed becomes wider with increasing strength of the radiation-driven shocks. Incidentally, these PDFs are consistent with those reported recently for star-forming clouds exposed to strong fluxes of {\\small IR}.  Multiple stellar-systems appearing in isolated pockets of gas compete for gas as envisaged in the standard competitive-accretion scenario. The protostellar {\\small MF}s derived in the present work, though fitted by a power-law, are not coeval. Interestingly, the {\\small MF}s for the irradiated cloud in cases 2, 3 and 4 better agree with the canonical {\\small IMF}, or even its more recent variation suggested by Kroupa. In fact, in case 4 where the radiation-induced shocks were the strongest, the {\\small MF} even demonstrated a turnover at $\\sim$0.2 M$_{\\odot}$ which is consistent with the Jeans mass for the shocked gas maintained at its pre-shock temperature. The {\\small MF} in case 1, though a power-law, is considerably shallow.Thus we have shown that cloud models with injected turbulence can possibly reconcile some key dynamical features of typical star-forming regions. This study, for instance, also shows that star-formation once triggered, proceeds rapidly and newer stars form within just over a few 10$^{4}$ years, which is significantly smaller than the free-fall time of a typical star-forming cloud. Also, the star-formation rate ({\\small SFR}) derived for the irradiated cloud in cases 2, 3 and 4 is consistent with that reported for typical star-forming clouds. Thus, despite the dissipative nature of turbulence, it can profoundly affect the star-formation history of a {\\small MC} and external sources such as a flux of {\\small IR} can continuously replenish the turbulent energy within that cloud. Consequently, principle sources of turbulence such as  proto-/stellar feedback in nearby star-forming clouds, or nuclear star-formation activity in galaxies must be poignant for controlling the global star-formation history.  \\begin{table} \\centering \\begin{minipage}{80mm} \\caption{Physical parameters for the mass function (MF) derived in each test case. The indices $\\alpha_{1}$ and $\\alpha_{2}$ are the respective slopes on either side of the knee (i.e., M$_{sink}\\ <$ M$_{knee}$ and  M$_{sink}\\ >$ M$_{knee}$) of the MF. The MF only in case 4 turnsover for (M$_{sink}$/M$_{\\odot})\\lesssim 0.2$, and the corresponding slope is denoted by $\\alpha_{3}$.} \\begin{tabular}{lccccc} \\hline \\hline Case & $\\alpha_{1}$ & $\\alpha_{2}$ & knee &  $\\alpha_{3}$ & turn-over \\\\ & & &   [M$_{\\odot}$] & & [M$_{\\odot}$] \\\\ \\hline 1  & 0.9 & 1.6 & 0.4  &MF incomplete \\\\ &&&& below $\\sim$0.2 M$_{\\odot}$\\\\ \\hline 2   & 1.3 & 2.0 & 0.2  & as above\\\\ \\hline 3  & 1.6 & 2.2 & 0.2 & as above\\\\ \\hline 4  & 2.0 & 2.4 & 0.4 & -1.3 & 0.2 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}"
    },
    "1208/1208.4111_arXiv.txt": {
        "abstract": "We use the progenitor of SN~2012aw to illustrate the consequences of modeling circumstellar dust using Galactic (interstellar) extinction laws that (1) ignore dust emission in the near-IR and beyond; (2) average over dust compositions, and (3) mis-characterize the optical/UV absorption by assuming that scattered photons are lost to the observer. The primary consequences for the progenitor of SN~2012aw are that both the luminosity and the absorption are significantly over-estimated.  In particular, the stellar luminosity is most likely in the range $10^{4.8} < L_*/L_\\odot < 10^{5.0}$ and the star was not extremely massive for a Type~IIP progenitor, with $M_* < 15M_\\odot$.  Given the properties of the circumstellar dust and the early X-ray/radio detections of SN~2012aw, the star was probably obscured by an on-going wind with $\\dot{M}\\sim 10^{-5.5}$ to $10^{-5.0} M_\\odot$/year at the time of the explosion, roughly consistent with the expected mass loss rates for a star of its temperature ($T_* \\simeq 3600_{-200}^{+300}$~K) and luminosity.  In the spirit of Galactic extinction laws, we supply simple interpolation formulas for circumstellar extinction by dusty graphitic and silicate shells as a function of wavelength ($\\lambda \\geq 0.3\\mu$m) and total (absorption plus scattering) V-band optical depth ($\\tau_V \\leq 20$).  These do not include the contributions of dust emission, but provide a simple, physical alternative to incorrectly using interstellar extinction laws. ",
        "introduction": "\\label{sec:introduction}  A key component to understanding supernovae (SNe) is the mapping between the explosions and their progenitor stars.  Slow, steady progress is being made, and it is now well established that Type~IIP SN are associated with red supergiants (see the review by \\citealt{Smartt2009}).  There is a puzzle, however, in that the observed upper limit of $\\sim 16 M_\\odot$ on the masses Type~IIP SN progenitors appears to be significantly lower than the maximum masses of $\\sim 25M_\\odot$ for  stars expected to explode while still red supergiants (\\citealt{Smartt2009b}), part of a more general absence of massive SN progenitors (see \\citealt{Kochanek2008}). Since these ``missing'' progenitors should be more luminous than those which are being discovered, they must either be hidden, evolve differently than expected, or fail to explode.  For example, in the rotating models of \\cite{Ekstrom2012}, the upper mass limit for red supergiants is lower, with $20M_\\odot$ stars being blue rather than red at the onset of carbon burning.  Alternatively, \\cite{Oconnor2011} and \\cite{Ugliano2012} find that the progenitor masses of $20$-$25M_\\odot$ corresponding to the upper mass range for red supergiants are more prone to failed explosions and prompt black hole formation, and such events would have to be found by searching for stars disappearing rather than explosions appearing (\\citealt{Kochanek2008}).  Binary evolution can also alter the distribution of final states at the time of explosion given the high probability for mass transfer as stars expand to become red supergiants (e.g. \\citealt{Sana2012}).  If the discrepancy is not explained by the physics of stellar evolution or explosion, then the most likely remaining explanation is the effect of circumstellar dust.  For example, \\cite{Walmswell2012} note that more massive and luminous red supergiants have stronger winds which can form dust and partly obscure the star.  This then biases (in particular) the upper mass limits associated with failed searches for progenitor stars, since the luminosity limits must be corrected for an unknown amount of circumstellar dust extinction.  \\cite{Fraser2012} and \\cite{Vandyk2012} recently analyzed the progenitor of SN~2012aw, finding that it was both relatively high mass ($15$-$20M_\\odot$ for \\cite{Vandyk2012} and $14$-$26M_\\odot$ for \\cite{Fraser2012}) and the most heavily obscured of any SN progenitor other than the completely obscured (and debated) SN~2008S class (see \\citealt{Prieto2008}, \\citealt{Thompson2009}, \\citealt{Kochanek2011}).  Since most of the extinction vanished after the SN, the dust must have been circumstellar (\\citealt{Fraser2012}, \\citealt{Vandyk2012}).  Like most studies of circumstellar dust in supernovae or supernovae progenitors, \\cite{Walmswell2012}, \\cite{Fraser2012} and \\cite{Vandyk2012} treat circumstellar dust as if it is a foreground screen that can be quantitatively modeled using  Galactic interstellar extinction curve models parametrized by the value of $R_V$ (e.g. \\citealt{Cardelli1989}). Usually only studies that are self-consistently calculating emission by the dust correctly model the absorption by circumstellar dust (e.g. studies of the SN~2008S class of transients by \\citealt{Wesson2010}, \\citealt{Kochanek2011}, or \\citealt{Szczygiel2012}).  But three well-known effects mean that it is never appropriate to make this approximation unless the optical depth is negligible compared to the required precision of the analysis.    First, emission from circumstellar dust can be important in the near-IR if the star is presently forming dust at temperatures of $1000$-$2000$~K.  This matters most for hotter stars with less intrinsic near-IR emission than the red supergiants we consider here.  Second, interstellar dust has the average composition of dust from all sources, while individual stars have the dust associated with their particular chemistry. This is relevant here because $\\sim 20 M_\\odot$ stars generally produce silicate rather than graphitic dusts (e.g. \\citealt{Verhoelst2009}).   \\begin{figure*} \\plotone{fig1.ps} \\caption{ DUSTY model of a $10^4$~K black body surrounded by $\\tau_V=3$ of cold silicate dust in a shell with $R_{out}/R_{in}=2$ and $\\rho \\propto 1/r^2$ in the shell.  The dust converts the input spectral energy distribution (SED) to the output SED, where the output SED is comprised of contributions from photons that escape without scattering (direct light) and photons that escape after being scattered (scattered light).  In an image, the direct emission is a point source and the scattered emission is a halo with a radius comparable to the inner radius of the shell. If the dusty shell is large enough to be resolved, only the direct emission is measured as coming from the source. We have made the dust cold enough to have no contributions from dust emission over this wavelength range. } \\label{fig:modsed} \\end{figure*}  \\begin{figure*} \\plotone{fig2.ps} \\caption{ Effective extinction laws $R_\\lambda$ if the dust is a foreground screen where we only measure the unscattered direct light.  The solid curves show the results for $\\tau_V=3$ of graphitic dust (top), silicate dust (bottom) and a 50:50 mix (middle). The dashed curves show the \\cite{Cardelli1989} extinction laws (CCM) with $R_V=4$ (top), $3.1$ (middle) and $2.0$ (bottom).  Normal Galactic dust ($R_V=3.1$) is relatively well modeled by the 50:50 mix. } \\label{fig:moddir} \\end{figure*}   \\begin{figure*} \\plotone{fig3.ps} \\caption{ Effective extinction laws $R_\\lambda$ for dust in an unresolved shell around the star where we measure both the direct and scattered light. The solid curves show the results for $\\tau_V=3$ of graphitic dust (top), silicate dust (bottom) and a 50:50 mix (middle). The dashed curves show the best fit CCM (\\citealt{Cardelli1989}) extinction laws with $R_V=2.4$ (top, graphitic), $R_V=2.1$ (middle, 50:50 mix) and $R_V=1.6$ (bottom, silicate). The fits are never very good, particularly at high optical depth, and the best fit parameters vary somewhat with the optical depth and wavelength. \\cite{Goobar2008} obtained similar results.  The peak at long wavelengths (small $\\lambda^{-1}$) is a silicate emission feature which DUSTY includes in the scattered light contribution to the emerging spectrum. } \\label{fig:modall} \\end{figure*}  The third, and least appreciated point, is the different role of scattered photons in interstellar and circumstellar extinction. This issue has been discussed in detail by \\cite{Wang2005} and \\cite{Goobar2008} in the context of anomalously low estimates of $R_V$ for Type~Ia SN, but correctly treating the problem has not become a matter of practice.  For a foreground screen, scattered light forms a very diffuse, extended halo around the source which is not included in the estimate of the source flux. This halo can sometimes be seen as the extended dust echoes of transient sources (e.g. \\citealt{Sugerman2002}).  If the circumstellar dust is unresolved, however, the scattered light is simply included in the total flux.  Thus, for interstellar extinction you observe only the direct, unabsorbed and unscattered emission, while for circumstellar extinction you observe both the direct unabsorbed and the scattered emission. {\\it Most of the observed optical emission from a moderately self-obscured star is scattered light.}  Fig.~\\ref{fig:modsed} shows an example generated by DUSTY (\\citealt{Ivezic1997}, \\citealt{Ivezic1999}, \\citealt{Elitzur2001}) using the spectral energy distribution (SED) of a $10^4$~K black body surrounded by $\\tau_V=3$ (scattering plus absorption) of cold silicate dust.  The dusty material has a density profile of $\\rho \\propto 1/r^2$ in a shell with $R_{out}/R_{in}=2$.  Galactic (interstellar) extinction corresponds to putting the dust at such a large radius that we no longer include the scattered light in the observed flux of the star, and the extinction law corresponds to the wavelength dependent difference between the input spectrum and the directly escaping spectrum.  Fig.~\\ref{fig:moddir} shows this ratio converted into an extinction law $R_\\lambda$ for pure graphitic, pure silicate and a 50:50 mix of the two, as compared to \\cite{Cardelli1989} Galactic extinction laws with $R_V=4.0$, $3.1$ and $2.0$.  Galactic dust models typically have a roughly 50:50 mix of graphitic and silicate grains (see the summary in \\citealt{Draine2011}), and using this mix in the DUSTY models roughly reproduces a typical $R_V=3.1$ Galactic extinction law.  The match is not perfect because there are differences in the assumed size distributions.  For a star surrounded by an unresolved shell of dust, however, the emission we measure is the sum of the direct and scattered light, which is very different from the direct emission alone.  Fig.~\\ref{fig:modall} shows this case converted into an effective circumstellar extinction law for the same three dust mixtures.  For circumstellar dust, all three examples are now significantly below the $R_V=3.1$ curve and a given change in color implies a significantly smaller change in luminosity. If we fit the absorption in these DUSTY models with \\cite{Cardelli1989} extinction models, the best fits for the graphitic, silicate and 50:50 dusts are $R_V\\simeq 2.4$, $R_V\\simeq 1.6$ and $R_V \\simeq 2.1$, respectively. These are the best fits for $\\tau_V=3$ from $0.36\\mu$m (U band) to $1\\mu$m.  The best fits vary modestly ($\\Delta R_V \\simeq 0.1$) with $\\tau_V$ and the fitted wavelength range, become worse for higher optical depths, and there are no truly good fits. Fig.~\\ref{fig:modall} superposes these \\cite{Cardelli1989} extinction curves on the DUSTY models.  \\cite{Goobar2008} found similar estimates for the best \\cite{Cardelli1989} extinction curves to model circumstellar dust and also noted that they are not very good approximations.  The final point to note in Figs.~\\ref{fig:moddir} and \\ref{fig:modall} is that the dust composition is quantitatively important.  Interstellar dust is a mixture of graphitic and silicate dusts from many sources. Individual stars, however, typically only produce silicate or graphitic dusts depending on the carbon to oxygen abundance ratio of the stellar atmosphere.  To zeroth order, all available carbon and oxygen bond to make CO. Then, a carbon poor (rich) atmosphere has excess oxygen (carbon) to make silicate (graphitic) dusts.  The atmospheres of $\\sim 20M_\\odot$ red supergiants usually form silicate dusts (e.g. \\citealt{Verhoelst2009}) because they never undergo the dredge up phases that enrich the atmospheres of the lower mass AGB stars with carbon (e.g. \\citealt{Iben1983}). Assuming a mixed interstellar composition will generally overestimate the absorption associated with a given change in color for silicate dusts and underestimate it for graphitic dusts.   In this paper we consider the effects of dust on the progenitor of SN~2012aw in more detail.  First, in \\S\\ref{sec:midir} we use archival Spitzer IRAC data to set stringent limits on the 3.6, 4.5, 5.8 and $8.0\\mu$m fluxes of the progenitor, and then model the spectral energy distribution (SED) of the progenitor with DUSTY in order to correctly model absorption, scattering and emission from a dust enshrouded star.  We highlight where any extinction law derived for dust screens at great distances from the source such as the standard \\cite{Cardelli1989} extinction laws can create problems. Finally, in \\S\\ref{sec:discussion} we summarize the results and provide simple interpolation formulas for extinction by graphitic and silicate circumstellar dust as a function of wavelength and optical depth.  \\begin{figure*} \\plotone{fig4.ps} \\caption{ Combined $3.6$ and $4.5\\mu$m images of the region surrounding the progenitor of SN~2012aw.  The 2\\farcs0 diameter circle marks the estimated position of the progenitor, although the formal uncertainties in the position are far smaller ($0\\farcs05$).  The smaller points show the grid of apertures used to estimate the flux limits, where missing points show the locations of dropped apertures (see text). A few of the grid points lie near fainter sources that are obvious in this co-added image but only marginally detectable (at about $2\\sigma$) in the four individual images combined to make Fig.~\\ref{fig:image}. The image is aligned to the world coordinate system with North up and East left. } \\label{fig:image} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion}  \\cite{Fraser2012} and \\cite{Vandyk2012} argue that the progenitor was probably the most massive yet found for a Type~IIP SN, probably at or above the upper limit of $(16.5\\pm1.5)M_\\odot$ \\cite{Smartt2009b} found in their statistical analysis of the masses of Type~IIP progenitors.  Given the amount of circumstellar extinction, this seemed to match the suggestion by \\cite{Walmswell2012} that the upper mass limit could be biased by increasing levels of dust formation for the more massive red supergiants due to the increase in mass loss rates with luminosity. Here we argue that many of these inferences are biased by incorrectly modeling circumstellar dust with an interstellar extinction law, thereby overestimating both the amount of extinction and the luminosity of the star.  When we model the SED using circumstellar dust models, the luminosity of the star is $L_* < 10^5 L_\\odot$ and the mass is $M_* < 15M_\\odot$, where the downwards shifts are easily understood from the differing physics of interstellar and circumstellar extinction. The visual extinction is overestimated because interstellar extinction neglects the contributions of scattered light,   and, to a lesser extent because of the low stellar temperature, the near-IR (K$_s$) extinction is overestimated by neglecting the emission from hot dust.  The absence of a mid-IR source noted by \\cite{Fraser2012} is a key point of evidence, since there should have been a detectable source given the proposed, higher luminosities unless the dust is cold and emitting at longer wavelengths than the IRAC bands.  In this particular case, the dust composition has little effect on the inferred properties of the star.  Unfortunately, without measuring the mid-IR portion of the SED we cannot determine the dust temperature $T_d$ other than limiting it to be lower than the dust destruction temperature ($T_d \\simeq 1500$~K). However, the wind properties are a strong function of the dust temperature because at fixed optical depth far more mass is required if the material is at larger radii and colder temperatures.  Ignoring the minor (10\\%) corrections from the finite value of $R_{out}/R_{in}=10$, the mass loss rate required to support the optical depth is \\begin{equation} \\dot{M} =  4 \\pi R_{in} v_w \\tau_V \\kappa_V^{-1} \\simeq 10^{-5} R_{in15} v_{w10} \\tau_{V5} \\kappa_{V100}^{-1} M_\\odot~\\hbox{year}^{-1} \\label{eqn:mdot} \\end{equation} where $R_{in} = 10^{15} R_{in15}$~cm, $v_w=10 v_{w10}$~km/s, $\\tau_V = 5\\tau_{V5}$ and $\\kappa_V = 10^2 \\kappa_{v100}$~cm$^2$/g. Table~\\ref{tab:sedmodels} gives the values for the individual models.  Supporting $\\tau_V \\sim 5$ when $T_d \\simeq 100$~K requires mass loss rates of order $\\dot{M} \\sim 10^{-2.5}M_\\odot$/year and implies total wind masses of order $3 M_\\odot$ that are implausible for a star with an initial mass of $M_*<15M_\\odot$. For comparison, the empirical approximations of \\cite{Dejager1988} predict $\\log_{10} (\\dot{M}/M_\\odot/\\hbox{year}) \\simeq -5.8 \\pm 0.5$ for a red supergiant wind, roughly matching the rates needed if the wind is producing dust at the time of the explosion.  The density of the wind is also tied to the expected phenomenology of the explosion.  If we assume the standard $\\rho_e \\propto v^{-12}$ outer ejecta density profile for red supergiants (\\citealt{Matzner1999}) then the expected shock velocity is \\begin{equation} v_s = 8200 E_{51}^{9/20} M_{e10}^{-7/20} \\dot{M}_{-4}^{-1/10} v_{w10}^{1/10} t_1^{-1/10}~\\hbox{km/s} \\end{equation} (e.g. \\citealt{Chevalier2003}), where the total energy of the supernova is $E= 10^{51} E_{51}$~ergs, the ejected mass is $M_e = 10 M_{e10}M_\\odot$, $\\dot{M}=10^{-4} \\dot{M}_{-4}M_\\odot$/year and $t_1$ is the elapsed time in days.  A shock expanding through a dense medium generates a luminosity of $L_S = (1/2)\\dot{M} v_s^3/v_w$, which we report in Table~\\ref{tab:sedmodels} for $v_s=5000$~km/s and $v_w=10$~km/s.  The emission from the forward shock is usually too hard to be easily detected, but assuming that the reverse shock is cooling and that its softer emissions dominate the observable X-ray emissions, the expected X-ray luminosity is \\begin{equation} L_x \\simeq { 9 \\dot{M} v_s^3 \\over 500 v_w} \\simeq 1.63 \\times 10^7 E_{51}^{27/20} M_{e10}^{-21/20} \\dot{M}_{-4}^{7/10} v_{w10}^{-7/10} t_1^{-3/10} L_\\odot \\label{eqn:xlum} \\end{equation} with a temperature of order 1~keV (e.g. \\citealt{Chevalier2003}). If we model the X-ray fluxes in Table~\\ref{tab:xray} using Eqn.~\\ref{eqn:xlum} and a range of additional absorption from $N_H=10^{20}$ to $10^{23}$~cm$^{-2}$, the mass loss rates for epochs 1, 3 and 5 (which have the smallest uncertainties) are \\begin{equation} \\dot{M} \\sim \\left( 10^{-6.4\\pm0.4}, 10^{-5.8\\pm0.3}, 10^{-4.7\\pm 0.3}\\right) v_{w10} M_{e10}^{3/2} E_{51}^{-27/14} M_\\odot/\\hbox{year}, \\end{equation} respectively, with some evidence that the amount of excess absorption needed above Galactic is increasing with time.   Such fluctuations and trends are not unusual (e.g. \\citealt{Dwarkadas2012}), but the X-ray emission appears to be broadly consistent with the presence of a wind with roughly the right density to explain the extinction of the progenitor. \\cite{Stockdale2012} and \\cite{Yadav2012} report rising $\\simeq 20$~GHz radio fluxes of $0.160\\pm0.025$  and $0.315\\pm0.018$~mJy roughly 7 and 13 days after discovery that also argue for a significant wind at the time of the SN. If we model the radio emission following  \\cite{Soderberg2005} assuming the ejecta mass is $M_e=(15-1.4)M_\\odot=13.6 M_\\odot$ and $E_{51}=1$, we obtain estimates of $\\dot{M} \\sim 10^{-5.0} v_{w10} M_\\odot$/year, although given only two data points at essentially the same frequency, the models are not tightly constrained.  The Thompson optical depth of the wind is always negligible, since it is a small ($\\ltorder 1\\%$) fraction of the dust optical depth, but the cold dust solutions would likely convert the SN into a Type~IIn because the H$\\alpha$ luminosity from recombination is of order $3000R_{in15}L_\\odot$ and increases linearly with the distance to the circumstellar material.  Thus, while the data is fragmentary, the simplest interpretation appears to be that there was a relatively steady $\\dot{M} \\sim 10^{-5.5}$ to $10^{-5.0}M_\\odot$/year wind creating the obscuration at the time of the SN.  Since the primary reason for inappropriately using Galactic extinction laws for circumstellar dust is almost certainly their ease of use, we supply in Table~\\ref{tab:models} equivalently easy to use models for absorption by circumstellar dust.   The problem is somewhat more complex because the extinction depends on both wavelength and optical depth, but the absorption in the DUSTY models from the UV to mid/near-IR ($0.3\\mu$m to $5.0\\mu$m) can be well modeled by the functional form \\begin{equation} A_\\lambda (\\tau_V) = \\tau_V \\lambda^{-x} \\sum_i \\sum_j a_{ij} \\tau_V^i \\lambda^{-j}. \\end{equation} for optical depths up to $\\tau_V=20$.  Here $\\lambda$ is the wavelength in microns and $\\tau_V$ is the total (absorption plus scattering) optical depth in the V band.  Table~\\ref{tab:models} provides these models for $R_{out}/R_{in}=2$ and $10$ in a format where they can simply be extracted from the electronic paper using a mouse and inserted into most numerical environments. The fits reproduce the DUSTY results with rms fractional residuals of 1.4\\% (1.6\\%) and 1.8\\% (2.2\\%) for the graphitic and silicate models and $R_{out}/R_{in}=10$ ($2$), although this includes some numerical rounding errors in the DUSTY models at low optical depths and longer wavelengths.  They can be extrapolated to longer wavelengths relatively safely, but at these longer wavelengths one would almost always also need to include dust emission.  {\\it They should not be extrapolated to shorter wavelengths or higher optical depths.} Fig.~\\ref{fig:opdepth} shows the circumstellar extinction of the DUSTY models used to build these interpolating functions and the interpolating functions as extracted from Table~\\ref{tab:models} for $R_{out}/R_{in}=10$.  In some circumstances it will also be useful to separate the direct and scattered light.  The fraction of the observed flux that is direct, $F_\\lambda^{abs}$, can be well-modeled by \\begin{equation} G_\\lambda^{abs} = -2.5 \\log_{10} F_\\lambda^{abs} (\\tau_V) = \\tau_V^{1/2} (1+\\lambda^{x})^{-1} \\sum_i \\sum_j a_{ij} \\tau_V^{i/2} \\lambda^{-j}. \\label{eqn:direct} \\end{equation} These interpolating functions are supplied in Table~\\ref{tab:direct} and the quality of the fits is shown in Fig.~\\ref{fig:direct}. For a source with intrinsic spectrum $S_\\lambda$, the observed spectrum is $S_\\lambda 10^{-0.4 A_\\lambda}$, the directly escaping flux is $S_\\lambda 10^{-0.4 A_\\lambda} F_\\lambda^{abs} = S_\\lambda 10^{-0.4(A_\\lambda+G_\\lambda^{abs})}$ and the escaping but scattered flux is $S_\\lambda 10^{-0.4 A_\\lambda} (1-F_\\lambda^{abs})$. Obviously, this approach can be generalized to other dust compositions or size distributions.  We parametrize the models by $\\tau_V$ because $\\tau_V$ is closely related to the physical properties of the wind (Eqn.~\\ref{eqn:mdot}), while relating it to a color (i.e. $E(B-V)$) would divorce the model from the underlying physics.  There is no comparably simple means of treating dust emission because it depends critically on the input spectrum.  These differences are quantitative rather than qualitative -- in circumstellar dust models the progenitor of SN~2012aw is still fairly heavily obscured, as found by \\cite{Fraser2012} and \\cite{Vandyk2012}, and \\cite{Walmswell2012} are correct that circumstellar dust introduces a bias that must be modeled when considering the statistics of SN progenitors, particularly when modeling non-detections in analyses like \\cite{Smartt2009b}.   However, since changes in extinction physics exponentially modify quantities like luminosities, these differences between the two dust geometries are quantitatively important.  We note, however, that SN where circumstellar dust will strongly bias inferences about the progenitor have densities of circumstellar material that will lead to X-ray or radio emission, as observed for SN~2012aw, because the optical depth is proportional to the wind density.  For example, combining Eqns.~\\ref{eqn:mdot} and  \\ref{eqn:xlum}, we see that the expected X-ray luminosity is $L_X \\propto \\tau_V$.  This means that the properties of the explosion can be used to constrain biases from dust around the progenitor even if the observations of the progenitor are inadequate to constrain the circumstellar extinction.  There are, of course, additional complexities coming from the geometry of the dust around the star that can lead to differences in the effective optical depth along the line of sight for direct emission and averaged over a significant fraction of the shell for the scattered emission.  Our simple models provide two means of approximating some of these effects.  First, changes in the shell thickness can be used to adjust the balance between scattered and absorbed light. Second, the emergent flux can be modeled as a sum of direct light with one optical depth, and scattered light with another in order to model differences between the mean and line-of-sight optical depths. {\\it Independent of these questions, interstellar and circumstellar extinction are quantitatively different, and in this case ignoring the differences leads to a significant overestimate of the progenitor luminosity and mass. }"
    },
    "1208/1208.5737_arXiv.txt": {
        "abstract": "CoGeNT employs p-type point-contact (PPC) germanium detectors to search for Weakly Interacting Massive Particles (WIMPs). By virtue of its low energy threshold and ability to reject surface backgrounds, this type of device allows an emphasis on low-mass dark matter candidates ($m_{\\chi}\\sim10$~GeV/c$^{2}$).  We report on the characteristics of the PPC detector presently taking data at the Soudan Underground Laboratory, elaborating on aspects of shielding, data acquisition, instrumental stability, data analysis, and background estimation.  A detailed background model is used to investigate the low energy excess of events previously reported, and to assess the possibility of temporal modulations in the low-energy event rate.  Extensive simulations of all presently known backgrounds do not provide a viable background explanation for the excess of low-energy events in the CoGeNT data, or the previously observed temporal variation in the event rate.  Also reported on for the first time is a determination of the surface (slow pulse rise time) event contamination in the data as a function of energy.  We conclude that the CoGeNT detector technology is well suited to search for the annual modulation signature expected from dark matter particle interactions in the region of WIMP mass and coupling favored by the DAMA/LIBRA results. ",
        "introduction": "CoGeNT (Coherent Germanium Neutrino Technology) is a program aiming to exploit the characteristics of p-type point-contact germanium detectors in areas as diverse as the search for low-mass dark matter candidates, coherent neutrino-nucleus elastic scattering, and $^{76}$Ge double-beta decay \\cite{jcap}.  Data collected from a first CoGeNT detector at a shallow underground location demonstrated sensitivity to low-mass ($<10$~GeV/c$^{2}$) dark matter particles \\cite{Aal08}. In particular, it appeared CoGeNT was particularly well suited to address the DAMA/LIBRA \\cite{DAMA} modulation result. Following the identification of several sources of internal background in this prototype, a second CoGeNT detector was installed in the Soudan Underground Laboratory (SUL) during 2009 with the goal of improving upon the dark matter sensitivity reach of the 2008 result \\cite{Aal08}. The first 56-days of operation of the CoGeNT detector at SUL showed an unexpected excess of events \\cite{Aal11} above the anticipated backgrounds for ionization energies below 2~keV. Further data collection from this detector continued until an interruption imposed by a fire in the access shaft to the laboratory halted the initial run in March of 2011. Analysis of the accumulated data set \\cite{Aal11b}, spanning 442 live days over the period 4 December 2009 to 6 March 2011, showed a $\\sim2.8\\sigma$ significance modulation of the monthly event rate in the low-energy region that is compatible with the dark matter signature described in \\cite{andrzej}. The fitting procedure generating this low-significance modulation result used unconstrained phase, period, and amplitude variables. Time-stamped data have been made publicly available, allowing for a number of independent analyses and interpretations.  In this paper we provide a more in-depth description of the apparatus and data analysis, concentrating on aspects of instrument stability, data cuts, uncertainties, and background estimation. The data set employed for this discussion is the same as in \\cite{Aal11b}, and all energies are in keVee (keV electron equivalent, i.e., ionization energy), unless otherwise stated. Following the three-month outage resulting from the Soudan fire, this detector has taken data continuously, starting 7 June 2011. An additional body of data is to be released in the near future. The design and expectations for CoGeNT-4 (C-4), a planned expansion aiming at an increase in active mass by a factor of ten, featuring four large PPC detectors with a reduced energy threshold and lower background, are discussed in a separate publication \\cite{inprep}. ",
        "conclusions": "CoGeNT is the first detector technology specifically designed to look for WIMP candidates in the low mass range around 10 GeV/c$^{2}$, an area of particular interest in view of existing anomalies in other dark matter experiments, recent phenomenological work in particle physics, and possible signals using indirect detection methods \\cite{indirect}. However, investigation of the largely unexplored $\\sim$few keV recoil energy range brings along new challenges in the understanding of low-energy backgrounds. The experience accumulated during the ongoing CoGeNT data-taking at SUL demonstrates that PPC detectors have excellent properties of long-term stability, simplicity of design, and ease of operation. This makes them highly suitable in searches for the annual modulation signature expected from dark matter particles forming a galactic halo.  Besides their excellent energy resolution, low energy threshold and ability to reject surface backgrounds, PPCs compare well to other solid-state detectors under several criteria: a) the relative simplicity of CoGeNT's data analysis results in comparable irreducible spectra regardless of analysis pipeline, b) the response to nuclear recoils is satisfactorily understood, resulting in a reliable nuclear recoil energy scale, c) uninterrupted stable operation of PPC detectors can be expected over very long (several year) timescales.  We plan to continue improving this technology and our understanding of low-energy backgrounds within the framework of a CoGeNT expansion, the C-4 experiment \\cite{inprep}."
    },
    "1208/1208.2051_arXiv.txt": {
        "abstract": "Using the results of Pichardo \\et (2005,2008), we determine regions of dynamical stability where planets (or discs in general) could survive in stable orbits around binary stellar systems. We produce this study for 161 binary stars in the Solar neighborhood with known orbital parameters. Additionally, we constructed numerically the discs (invariant loops) around five binary systems with known orbital parameters and with confirmed planets: HIP 10138, HIP 4954, HIP 67275, HIP 116727 and Kepler 16, as a test to the approximation of Pichardo et al. (2005,2008). In each single case, the reported position of the planets lay within our calculated stability regions. This study intends to provide a guide in the search for planets around binary systems with well know orbital parameters, since our method defines precise limits for the stable regions, where discs may have established and planets formed. ",
        "introduction": "It is known that most low-mass main-sequence stars are members of binary or multiple systems (Duquennoy \\& Mayor 1991; Fisher \\& Marcy 1992), and in particular in the Solar Neighborhood, the fraction goes up to $\\sim 78\\%$ (Abt 1983). This suggests that binary formation is the primary branch of the star formation process (Mathieu 1994).  Significant advances in high-angular-resolution infrared imaging technology have enabled large surveys of young binary stars on a variety of star-forming regions (Mathieu \\et 1992, 1994). In addition, right after the discovery of the first extrasolar planetary system around a pulsar (Wolszczan \\& Frail 1992), and particularly after the first extrasolar planet discovered around a main sequence star (Mayor \\& Queloz 1995; Marcy \\& Butler 1998), observational activity was greatly stimulated. More recently, advances in observational techniques and instrumentation, such as the HST (WFPC2 \\& NICMOS) imaging (Padgett \\et 1997, 1999; Reid \\et 2001; Borucki \\et 2010), submillimeter imaging (Smith \\et 2000), optical and infrared long-baseline interferometry (Quirrenback 2001a,b), millimeter and submillimeter interferometry (Launhardt \\et 2000, Launhardt 2001, Guilloteau 2001), adaptive optics (Simon \\et 1999; Close 2001), spatial astrometry (S\\\"oderhjelm 1999; Quist \\& Lindegren 2000, 2001), and microlensing (Alcock \\et 2001; Dong-Wook \\et 2008; Rattenbury 2009), are also available in binary studies. Thanks to all this new technology and observational work, we have now the possibility to study and understand better the physics of binary systems and the surrounding discs built during the formation stage.  In the recent past, several planets in binary or multiple star systems have been discovered (Correia 2008; Deeg \\et 2008; Desidera \\& Barbieri 2007; Fischer \\et 2008; Raghavan 2006; Konacki 2005; Mugrauer \\et 2005; Eggenberger \\et 2004; Sigurdsson \\et 2003; Udry \\et 2002; Sigurdsson \\& Phinney 1993; Lyne 1988, etc.). Both suspected of formed in {\\it situ}, or acquired by dynamical processes (Pfahl \\& Muterspaugh 2006). For a review of observational techniques see Muterspaugh \\et (2010). In addition, several recently discovered circumstellar discs where planets are assumed to be formed, lie around close binaries (Wright \\et 2011; Prato \\& Weinberger 2010; Desidera \\& Barbieri 2007; Doyle \\et 2011; Queloz \\et 2000; Hatzes \\et 2003). In these cases the presence of a companion star should have a very strong influence on both discs and planets. From the several hundreds of extrasolar planets confirmed so far (see http://exoplanet.eu, http://planetquest.jpl.nasa.gov, www.exoplanets.org) to be around main sequence stars, about 10\\% are known to reside in binary systems with a wide range of orbital separations. In almost all cases, the planet orbits in S-type configurations (Dvorak 1986), while the second star acts as a perturber to the planetary system. A circumbinary planet (P-type orbit) has recently also been detected in Kepler 16 (Doyle et al. 2011). This has motivated the search for stable periodic orbits around binary systems where planets (and discs in general) can settle down in a stable configuration. Most theoretical studies have focused on binaries in near-circular orbits (H\\'enon 1970; Lubow \\& Shu 1975; Paczy\\'nski 1977; Papaloizou \\& Pringle 1977; Rudak \\& Paczy\\'nski 1981; Bonell \\& Bastien 1992; Bate 1997; Bate \\& Bonnell 1997). Even very precise analytical methods to approximate periodic orbits in circular binaries are available (Nagel \\& Pichardo 2008).  Due to the lack of conservation of the Jacobi integral, the case of eccentric binaries is qualitatively more complicated. Artymowicz \\& Lubow (1994) and Pichardo \\et (2005, 2008, hereafter PSA1 and PSA2) calculate the extent of zones in phase space available for stable, non self-intersecting orbits around each star and around the whole system. In this study we use the results of PSA1 and PSA2, to calculate stable regions for planets or discs in binary systems in the Solar Neighborhood with known orbital parameters such as, mass ratio, eccentricity and semimajor axis.  Although the existence of stable zones, as the ones we are calculating here, is a necessary, but not a sufficient condition to the existence of planets (or discs in general), if any stable material (planets, gas, etc.), exists in a given system, irrespective of their formation mechanism, they would be necessarily located within the limits of the stable zones. In this direction, a fruitful line of investigation is the intersection between phase space available zones for the long term evolution of planetary systems and habitable regions allowed by the binary system (Haghighipour \\et 2007; Haghighipour \\et 2010).  We present a table with the compilation of all the binaries in the Solar Neighborhood with known orbital parameters from different sources. In the same table are presented the results for our calculated circumstellar and circumbinary stable zones around each binary system.  This paper is organized as follows. In Section \\ref{method}, we explain briefly the method employed to calculate regions of stable non self-intersecting orbits around binary stars. The binary star sample is presented in Section \\ref{sample}. Results, including stable regions of orbital stability for circumstellar and circumbinary planetary discs (and discs in general), and an application to observations of five binaries with observed planets, are given in Section \\ref{results}. In Section \\ref{conclusions} we present our conclusions. ",
        "conclusions": "We have compiled a sample of binary stars with known orbital parameters (semimajor axes, eccentricities and stellar masses) of the Solar neighborhood and present some basic statistics.  We calculate on this binary stars sample the extent of regions of stable non-self intersecting orbits where planets may exist. For this purpose, we have applied the concept of ``invariant loops'' and used the formulae of PSA1 and PSA2. Our approximation is ballistic, thus, the application is straightforward to debris discs (planets, cometary nucleii, asteroid belts, etc.). In the case of gas discs, further physics may constraint discs sizes, however, what we are providing here are the regions where the most important orbits of the binary, i.e., the ones that represent the backbone of the dynamical system (the ones that are followed for the most of the orbits), lay. We have computed the spatial limits of these circumstellar and circumbinary zones for a sample of 161 binaries in the Solar neighborhood where orbital data is known and presented it in the form of a table where all the relevant parameters are provided.  We compare our results with observations in the 5 cases where planets have been discovered in binary systems, and where semimajor axis for the planets are provided. We find that all the planets lay down within our computed regions of stability. In particular, for HD 120136, our predicted region of circumstellar stability is very small, and yet the discovered planet lays within it. Although confrontation with a larger database is desirable, the current statistics is fully consistent with our results, proving reliable our approach. The tool of ``invariant loops'' may be very helpful in the search for planets in binary systems."
    },
    "1208/1208.0781_arXiv.txt": {
        "abstract": "\\noindent We have undertaken a new ground-based monitoring campaign on the broad-line radio galaxy 3C\\,390.3 to improve the measurement of the size of the broad emission-line region and to estimate the black hole mass. Optical spectra and {\\it g}-band images were observed in late 2005 for three months using the 2.4-m telescope at MDM Observatory. Integrated emission-line flux variations were measured for the hydrogen Balmer lines H$\\alpha$, H$\\beta $, H$\\gamma$, and for the helium line \\heii $\\lambda 4686$, as well as {\\it g}-band fluxes and the optical AGN continuum at $\\lambda =$\\,5100\\,\\AA . The {\\it g}-band fluxes and the optical AGN continuum vary simultaneously within the uncertainties, $\\tau _{cent} = (0.2\\pm1.1)$\\,days. We find that the emission-line variations are delayed with respect to the variable {\\it g}-band continuum by $\\tau (H\\alpha)        = 56.3^{+2.4}_{-6.6}$\\,days, $\\tau (H\\beta )         = 44.3^{+3.0}_{-3.3}$\\,days, $\\tau (H\\gamma)        = 58.1^{+4.3}_{-6.1}$\\,days, and $\\tau ($\\heii \\,4686$) = 22.3^{+6.5}_{-3.8}$\\,days. The blue and red peak in the double peaked line profiles, as well as the blue and red outer profile wings, vary simultaneously within $\\pm3$ days. This provides strong support for gravitationally bound orbital motion of the dominant part of the line emitting gas. Combining the time delay of the strong Balmer emission lines of H$\\alpha$ and H$\\beta$ and the separation of the blue and red peak in the broad double-peaked profiles in their rms spectra, we determine M$_{bh}^{vir} = 1.77^{+0.29}_{-0.31}\\,\\times\\,10^8 M_\\odot$ and using $\\sigma_{line}$ of the rms spectra M$_{bh}^{vir} = 2.60^{+0.23}_{-0.31}\\,\\times\\,10^8 M_\\odot$ for the central black hole of 3C\\,390.3, respectively. Using the inclination angle of the line emitting region which is measured from superluminal motion detected in the radio range, accretion disk models to fit the optical double-peaked emission line profiles, and X-ray observations, the mass of the black hole amounts to M$_{bh} = 0.86^{+0.19}_{-0.18}\\,\\times10^9 M_\\odot$ (peak-separation) and M$_{bh} = 1.26^{+0.21}_{-0.16}\\,\\times10^9 M_\\odot$ ($\\sigma_{line}$), respectively. This result is consistent with the black hole masses indicated by simple accretion disk models to describe the observed double-peaked profiles, derived from the stellar dynamics of 3C\\,390.3, and with the AGN radius\\,--\\,luminosity relation. Thus, 3C\\,390.3 as a radio-loud AGN with a low Eddington ratio, L$_{edd}$/L$_{bol}$ = 0.02, follows the same AGN radius\\,--\\,luminosity relation as radio-quiet AGN. ",
        "introduction": "The mass of super-massive black holes (SMBHs) is of fundamental importance for understanding active galactic nuclei (AGN) which are powered by mass accretion onto a SMBH. Analysis of AGN variability by applying the technique of reverberation mapping (RM) has been established as a powerful tool to determine the size of the broad emission-line region (BLR) of AGN (Blandford \\& McKee 1982; Horne et al.\\,2004; Netzer \\& Peterson 1997; Peterson 2003) and, under the assumption of virial gas motion, consequently the SMBH mass (e.g., Peterson et al.\\,2004).  Large observational efforts are required to obtain data of sufficient quality and temporal coverage to study AGN variability and to employ RM methods. Currently, about 50 AGN have been monitored for periods of at least a few months, usually as part of international collaborations. The entire available database of reverberation observations obtained through 2003 was uniformly re-analyzed with improved methods of time series analysis to minimize systematic errors (Peterson et al.\\,2004). Some of these results have been superseded by more recent experiments (Bentz et al.\\,2006b,\\,2007a; Denney et al.\\,2006,\\,2009a; Grier et al.\\,2008; Woo et al.\\,2010) and additional reverberation results steadily increase the size and quality of the database (e.g., Bentz et al.\\,2009b,\\,2010a; Denney et al.\\,2009b,\\,2010).  It has been shown that AGN and non-active galaxies follow the same relation between the black hole mass, M$_{bh}$, and stellar velocity dispersion, $\\sigma_\\ast$, of the central spheroidal component of the host galaxy (Ferrarese et al.\\,2001; Gebhardt et al.\\,2000; G\\\"{u}ltekin et al.\\,2009; Merritt \\& Ferrarese 2001; Tremaine et al.\\,2002) and that the masses based on RM studies are consistent with the masses derived from the M$_{bh}$ - $\\sigma_\\ast$ relation (Nelson et al.\\,2004; Onken et al.\\,2004; Woo et al.\\,2010). The close relations of the black hole mass with physical properties of the host galaxy, e.g. stellar velocity dispersion, bulge mass and bulge luminosity, indicate a coupled growth of the black hole and the formation and evolution of the host galaxy (e.g., Silk \\& Rees 1998; Haiman \\& Loeb 1998,2001; Bromm \\& Loeb 2003; Di\\,Matteo et al.\\,2004; Yoo \\& Miralda-Escud\\'e 2004; Volonteri \\& Rees 2006; Volonteri \\& Begelman 2010).  The broad-line radio galaxy (BLRG) 3C\\,390.3 is a bright and nearby ($m_v = 15.0$, $z=0.056$; Osterbrock, Koski, \\& Phillips 1976) FR\\,II radio galaxy with extended double-lobed radio emission (Leahy \\& Perley 1995). After the identification as optical counterpart of the radio source 3C\\,390.3 (Wyndham 1966), it was classified as a N-type galaxy by Sandage (1966). Soon, it was discovered that 3C\\,390.3 shows very broad Balmer emission-lines (Lynds 1968) which were later recognized as prominent double-peaked emission-line profiles (Burbidge \\& Burbidge 1971). These double-peaked profiles are generally regarded as a characteristic signature of accretion disk emission (e.g., Eracleous \\& Halpern 1994,2003; Gezari, Halpern, \\& Eracleous 2007).  So far almost all AGN variability studies have focused on radio-quiet AGN. However, 3C\\,390.3 has a well-known variability history (e.g., Cannon, Penston, \\& Penston 1968; Selmes, Tritton, \\& Wordsworth 1975; Barr et al.\\,1980; Penston \\& Perez 1984; Veilleux \\& Zheng 1991; Zheng 1996; Wamsteker et al.\\,1997; O'Brien et al.\\,1998; Dietrich et al.\\,1998; Sergeev et al.\\,2002; Tao et al.\\,2008), with photometric measurements going back to 1968 (Yee \\& Oke 1981). Using the Harvard plate collection Shen, Usher, \\& Barrett (1972) traced back brightness variations to 1895. Thus, 3C\\,390.3 was a prime target of a multiwavelength monitoring campaign in 1994/95 especially because of previous reports of dramatic changes in the Balmer line profile shape and strength and to perform coordinated X-ray/UV/optical monitoring of a radio-loud AGN for the first time. Currently, 3C\\,120 and 3C\\,390.3 are the only radio-loud AGN which have been monitored in detail. The multiwavelength study of 3C\\,390.3 in 1994/95 (Leighly et al.\\,1997; O'Brien et al.\\,1998; Dietrich et al.\\,1998) shows significant variations in the X-ray, UV, and optical domain over a period of one year. However, the light curves of the UV and optical variations show a nearly monotonic increase of the continuum and emission-line flux with only moderately strong substructures. The measured delays of several broad emission-lines relative to the observed continuum variations are about $\\tau \\simeq 20\\pm6$\\,days for H$\\beta $ and H$\\alpha$ (Dietrich et al.\\,1998) and $\\tau \\simeq 36$ to 60\\,days ($\\pm 18$\\,days) for \\civ\\ and Ly$\\alpha$ (O'Brien et al.\\,1998). More recently, Sergeev et al.\\,(2002) studied the variability characteristics of 3C\\,390.3 for a period of nearly one decade. They found a longer delay of the response of the  H$\\beta $ emission with respect to continuum variations compared to the result of the 1994/95 study. They concluded that the difference might be caused by different continuum variability characteristics which manifest in different continuum auto-correlation functions. Recently, Shapovalova et al.\\,(2010), Jovanovic et al.\\,(2010), Popovic et al.\\,(2011), and Sergeev et al.\\,(2011) discussed profile variations over periods from a few years up to about 15 years.  In the following work, we present the results of a monitoring campaign undertaken at MDM Observatory in late 2005. Based on imaging and spectroscopic data, we find that 3C\\,390.3 clearly showed broad band variations, delayed variations of the broad Balmer emission-lines H$\\alpha$, H$\\beta $, and H$\\gamma $, and of the helium line \\heii $\\lambda 4686$. Using the time delays and the peak separation in the rms spectra of these emission lines, we estimate a virial black hole mass of M$_{bh}^{vir} = 1.77^{+0.29}_{-0.31}\\,\\times\\,10^8 M_\\odot$ for 3C\\,390.3 and using $\\sigma_{line}$ of the rms spectra M$_{bh}^{vir} = 2.60^{+0.23}_{-0.31}\\,\\times\\,10^8 M_\\odot$, respectively. The detection of superluminal motion for 3C\\,390.3 indicates an inclination angle of $i = 27^o \\pm 2^o$ which allows to correct the observed velocity for the tilt of the line emitting region relative to the observer. Taking this into account, we derive a black hole mass of M$_{bh} = 0.86^{+0.19}_{-0.18}\\times 10^9 M_\\odot$ (peak separation) and M$_{bh} = 1.26^{+0.21}_{-0.16}\\times 10^9 M_\\odot$ ($\\sigma_{line}$) for 3C\\,390.3, respectively, which is, within the uncertainties, consistent with recent results based on the \\caii $\\lambda \\lambda 8494, 8542,8662$ stellar absorption triplet and the $M_{bh} - \\sigma_\\ast$ relation (Lewis \\& Eracleous 2006; Nelson et al.\\,2004). ",
        "conclusions": "\\subsection{Times Series Analysis and the Size of the BLR} Recently, results of studies on the long-term variability properties of 3C\\,390.3 have been presented by Sergeev et al.\\,(2002,\\,2011) and Shapovalova et al.\\,(2010). In two studies Sergeev et al.\\,(2002,\\,2011) investigated the correlated variations of the optical continuum and the response of the broad H$\\beta$ emission-line flux for the years 1992 to 2000 and 2000 to 2007, respectively. In both studies they found that the H$\\beta$ variations are delayed by $\\tau (H\\beta) = 82^{+12}_{-10}$\\,days ($1992 - 2000$) and $\\tau (H\\beta) = 94\\pm6$\\,days ($2000 - 2007$) and for H$\\alpha$ delays of $\\tau (H\\alpha) = 162^{+32}_{-15}$\\,days ($1992 - 2000$) and $\\tau (H\\alpha) = 174\\pm16$\\,days ($2000 - 2007$) were determined. A comparable result was found by Shapovalova et al.\\,(2010) who studied the optical variations for 3C\\,390.3 from 1995 to 2007. They report that the variations of the broad H$\\beta$ and H$\\alpha$ emission-line flux are delayed by $\\tau (H\\beta) = 96^{+28}_{-47}$\\,days and $\\tau (H\\alpha) = 127\\pm18$\\,days, both too long to be measurable from our data. However, due to the lower sampling rate of the H$\\alpha$ light curve the ICCF analysis displays two possible peaks also at $\\tau (H\\alpha) \\simeq 24$\\,days and $\\tau (H\\alpha) \\simeq 151$\\,days, respectively.  Those results are inconsistent with the delays which were reported by Dietrich et al.\\,(1998) using the data of the X-ray\\,--\\,UV\\,--\\,optical monitoring campaign in 1994/95 (Leighly et al.\\,1997; O'Brien et al.\\,1998; Dietrich et al.\\,1998), i.e., $\\tau (H\\beta) = 23\\pm4$\\,days and $\\tau (H\\alpha) = 19\\pm9$\\,days and also with the results of this study (Table 6). However, the time delays which we determined for 3C\\,390.3 in 1995 are consistent with the delays we find for H$\\alpha$ and H$\\beta$ in this study. In 1995 the strength of the optical continuum corrected for host-galaxy contributions was F$_\\lambda$(5100\\AA )\\,$= 1.16\\times10^{-15}$ erg\\,s$^{-1}$\\,cm$^{-2}$\\,\\AA $^{-1}$. During the 2005 monitoring campaign the continuum was about $\\sim 6\\times$ stronger (Table 5). Using the radius\\,--\\,luminosity relation for AGN (Bentz et al.\\,2009a) it can be expected that the derived delays of the H$\\alpha$ and H$\\beta$ emission-lines should be about $\\sim 2.5\\times$ as long than those of 1995. This is in good agreement with the delays which we have measured (Table 6).  To investigate the cause of the discrepancy between the results based on the analysis of the variations over about 10 years compared with those covering about three months up to one year, we re-analyzed the data published by Sergeev et al.\\,(2002,2011) and Shapovalova et al.\\,(2010). Using the data of the Sergeev et al. study, we compiled H$\\beta$ and F$_\\lambda$(5100\\AA ) continuum light curves which cover about 15 years with 413 epochs for the continuum and 131 epochs for H$\\beta$. We applied the same ICCF analysis to these light curves as we did for the measurements for this study. We found $\\tau (H\\beta) = 81.6^{+16.4}_{-22.5}$\\,days which is consistent with the results of Sergeev et al. and due to the large uncertainties, within 1-$\\sigma$ it is even in the range with the result of this study. Next, we studied parts of the light curve which covered approximately 1000 days to study the impact of the luminosity state of 3C\\,390.3 on the delay. We found delays of the broad H$\\beta$ variations ranging from $\\tau (H\\beta) = 13^{+34}_{-30}$\\,days to $\\tau (H\\beta) = 102^{+11}_{-16}$\\,days with no clear relation between the continuum state and the delay. However, the mean spacing, especially for the H$\\beta$ emission-line flux measurements, is only $\\sim20$\\,days to $\\sim160$\\,days, i.e., probably not sufficient. Only for the last three years of the light curve the temporal sampling is about $\\sim$$6$\\,days for the continuum and $\\sim$$20$\\,days for the H$\\beta$ emission-line flux. For this period we find a delay of $\\tau (H\\beta) = 102^{+6}_{-6}$\\,days. The H$\\beta$ emission line flux measurements by Sergeev et al.\\,(2011) and by Shapovalova et al.\\,(2010) do overlap with our monitoring campaign. Therefore, we can compare those with our H$\\beta$ emission line light curve (Fig.\\,18). It can be seen that the continuum light curve follows the variations we measured while the coverage in H$\\beta$ is not as high.  In the same way, we re-analyzed the continuum and H$\\beta$ emission-line light curves given by Shapovalova et al.\\,(2010). For the entire data set which covers about 12.5 years we find $\\tau (H\\beta) = 96.2^{+35.4}_{-35.2}$\\,days, consistent with the results of Sergeev et al. Similar to the Sergeev et al. data the light curves were split up into periods of about 1000 days duration with temporal sampling of about 40 days. We found delays of the broad H$\\beta$ variations ranging from $\\tau (H\\beta) = 16^{+39}_{-41}$\\,days to $\\tau (H\\beta) = 210^{+55}_{-140}$\\,days with no clear relation between the continuum state and the delay.  In a study on the reliability of cross-correlation function time delay determinations Welsh (1999) has pointed out that variations on longer time scales, e.g., on the dynamical time scale of the BLR which is of the order of years, will shift the response time to longer delays. Long term variations are different from direct reverberation signals which measure the instantaneous response of the emission-line gas to continuum variations, while long term variations trace a more gradual response to an overall increase or decrease of the continuum strength. To correct for this effect it is necessary to detrend light curves from those gradual changes. To detrend the continuum and H$\\beta$ emission-line variations, we fit a low order polynomial to the light curves. Next, we applied the analog time series analysis to the detrended light curves of the Sergeev et al.\\,(2002,2011) and Shapovalova et al.\\,(2010) studies. We derived time delays of the H$\\beta$ variations of $\\tau (H\\beta) = 90.4^{+8.9}_{-9.2}$\\,days (Sergeev et al.) and $\\tau (H\\beta) = 82.7^{+13.9}_{-12.9}$\\,days (Shapovalova et al.). Within the uncertainties (1-$\\sigma$\\,errors) these delays are consistent with those obtained with the original light curves. However, the uncertainties of the time delays using the detrented light curves are about a factor 2 to 3 smaller than those for the light curves including gradual long term variations.  Furthermore, the properties of continuum strength variations have a significant effect on the time delay measured from cross-correlation functions. This issue can be responsible for the different delays measured in our study and those using variations of the continuum and the broad emission lines over more than 10 years, in particular, the auto-correlation function of the continuum, ACF$_{cont}$. It has been already noted by Sergeev et al.\\,(2002,2011) that the width of the ACF$_{cont}$ of the continuum light curve which is covering several years up to more than a decade is much broader than the ACF$_{cont}$ of a shorter campaign like the one in 1994/95 and that this will result in a longer time delay (see also the Appendix for a more detailed discussion).  In addition, our campaign covered only a little more than 80 days and hence we are not able to measure time delays of about 100 days and more. We think that the different time delays are caused by the long time period of about a decade to cover the variation of 3C\\,390.3 and the wider temporal sampling of the measurements. Furthermore, the significantly different widths of the ACF$_{cont}$ and the fact that in addition to reverberation signals, i.e., the direct response of line emitting gas to continuum variations, also variations of the emission-line flux are included. These are associated with changes of the physical conditions and the distribution of the gas which happen on dynamical time scales and which are uncorrelated with continuum variability also contributes {\\bf to} the different time delays. Therefore, a long duration campaign for 3C\\,390.3 with densely sampled measurements will be necessary to find a definitive result.  \\subsection{Black Hole Mass Estimates} A wide range of mass estimates for the super-massive black hole of 3C\\,390.3 have been reported, ranging from $\\sim 1.3\\times10^8\\,M_\\odot$ up to $\\sim 7\\times10^9\\,M_\\odot$ using emission line profile properties and estimates of the size of the BLR (e.g., Barr et al.\\,1980; Bentz et al.\\,2009b; Clavel \\& Wamsteker 1987; Gaskell 1996; Peterson et al.\\,2004; Sergeev et al.\\,2002,2011; Wamsteker et al.\\,1997), as well as using the \\caii\\ triplet in the near infrared employing the M\\,--\\,$\\sigma_\\ast$ relation yielding about M$_{bh}\\simeq 4$ to $5\\times$ $10^8\\,M_\\odot$ (Nelson et al.\\,2004). Most of these studies favor a black hole mass for 3C\\,390.3 of the order of 5 to $10\\times$ $10^8\\,M_\\odot$. Our measurement of the mass of the black hole of 3C\\,390.3 with $M_{bh} = 0.86^{+0.19}_{-0.18}\\,\\times\\,10^9\\,M_\\odot$ (based on the separation of the blue and red peak in the rms spectrum) and $M_{bh} = 1.26^{+0.21}_{-0.16}\\,\\times\\,10^9\\,M_\\odot$ (using $\\sigma_{line}$) is consistent with the black hole mass based on the $M - \\sigma_\\ast $ relation. Using the M-$\\sigma_\\ast$ relation from G\\\"{u}ltekin et al.\\,(2009), we calculated the stellar velocity dispersion $\\sigma_\\ast$ which is expected for the host galaxy of 3C\\,390.3 based on our estimated black hole mass. We find $\\sigma_\\ast = 257^{+90}_{-70}$\\,km\\,s$^{-1}$ for 3C\\,390.3 which is consistent with $\\sigma_\\ast = 273\\pm16 $\\,km\\,s$^{-1}$ as measured by Nelson et al.\\,(2004). Furthermore, simple disk models which we applied to describe the overall structure of the broad double-peaked hydrogen Balmer emission lines yield a mass estimate of $M_{bh}\\simeq 10^9\\,M_\\odot$.  An additional test for the reliability of the derived black hole mass for 3C\\,390.3 is given by the comparison of our measured values for the size and continuum luminosity of 3C\\,390.3 with the expected values derived from the radius\\,--\\,luminosity relation. Guided by simple photoionization models, a relation between the continuum luminosity of an AGN and the radius of the BLR is expected and Kaspi et al.\\,(2000) provided the first convincing evidence for such a relation. Careful re-analysis and additional observations (e.g., Peterson et al.\\,2004; Bentz et al.\\,2006b,2007a,2009b; Denney et al.\\,2006,2009a,2010; Grier et al.\\,2008; Onken et al.\\,2003) have reduced the uncertainty of the slope of the relation. It turned out that the correction for host galaxy contamination has a profound impact on the R\\,--\\,L relation (Bentz et al.\\,2006a,2009a). With a slope of $\\alpha =0.52\\pm0.04$ (Bentz et al.\\,2009a), we estimated the black hole mass of 3C\\,390.3. For the continuum flux at $\\lambda = 5100$\\,\\AA\\ we used the average continuum flux of the AGN continuum (Table 5), corrected for host galaxy contributions, with $\\lambda\\,L_\\lambda(5100{\\rm \\AA}) = 2.28\\times10^{44}$ erg\\,s$^{-1}$ and the measured delay for the H$\\beta$ emission line is $\\tau _{cent} \\simeq 44$ days. Using the radius\\,--\\,luminosity relation as given in Bentz et al.\\,(2009a)  \\begin{equation} \\log\\,R_{BLR} = K + \\alpha \\,\\log\\,(\\lambda\\,F_\\lambda (5100{\\rm \\AA})) \\end{equation}  \\noindent with K\\,=\\,$-21.3^{+2.9}_{-2.8}$ and $\\alpha = 0.519^{+0.063}_{-0.066}$ the H$\\beta$ BLR radius amounts to $\\tau = 52.7^{+2.9}_{-2.8}$ days, which given the intrinsic scatter in the relationship, is consistent with the measured $\\tau _{cent}(H\\beta) = 44.3^{+3.0}_{-3.3}$\\,days or the $\\tau _{cent}(H\\beta) = 47.9^{+2.4}_{-4.2}$\\,days using SPEAR."
    },
    "1208/1208.1264_arXiv.txt": {
        "abstract": "A method is presented for automated photometric classification of supernovae (SNe) as Type-Ia or non-Ia. A two-step approach is adopted in which: (i) the SN lightcurve flux measurements in each observing filter are fitted separately to an analytical parameterised function that is sufficiently flexible to accommodate vitrually all types of SNe; and (ii) the fitted function parameters and their associated uncertainties, along with the number of flux measurements, the maximum-likelihood value of the fit and Bayesian evidence for the model, are used as the input feature vector to a classification neural network (NN) that outputs the probability that the SN under consideration is of Type-Ia.  The method is trained and tested using data released following the SuperNova Photometric Classification Challenge (SNPCC), consisting of lightcurves for 21,319 SNe in total. We consider several random divisions of the data into training and testing sets: for instance, for our sample ${\\cal D}_1$ (${\\cal D}_4$), a total of 10 (40) per cent of the data are involved in training the algorithm and the remainder used for blind testing of the resulting classifier; we make no selection cuts. Assigning a canonical threshold probability of $p_{\\rm th}=0.5$ on the network output to class a SN as Type-Ia, for the sample ${\\cal D}_1$ (${\\cal D}_4$) we obtain a completeness of 0.78 (0.82), purity of 0.77 (0.82), and SNPCC figure-of-merit of 0.41 (0.50).  Including the SN host-galaxy redshift and its uncertainty as additional inputs to the classification network results in a modest 5--10 per cent increase in these values. We find that the quality of the classification does not vary significantly with SN redshift. Moreover, our probabilistic classification method allows one to calculate the expected completeness, purity and figure-of-merit (or other measures of classification quality) as a function of the threshold probability $p_{\\rm th}$, without knowing the true classes of the SNe in the testing sample, as is the case in the classification of real SNe data. The method may thus be improved further by optimising $p_{\\rm th}$ and can easily be extended to divide non-Ia SNe into their different classes. ",
        "introduction": "\\label{sec:intro} Much interest in supernovae (SNe) over the last decade has been focused on Type-Ia (SNIa) for their use as `standardizable' candles in constraining cosmological models.  Indeed, observations of SNIa led to the discovery of the accelerated expansion of the universe (\\citealt{1998AJ....116.1009R}; \\citealt{1999ApJ...517..565P}), which is usually interpreted as evidence for the existence of an exotic dark energy component.  Ongoing observations of large samples of SNIa are being used to improve the measurement of luminosity distance as a function of redshift, and thereby constrain cosmological parameters further (e.g., \\citealt{2009ApJS..185...32K}; \\citealt{2012MNRAS.419..513B}; \\citealt{2011ApJ...737..102S}; \\citealt{2011ApJS..192....1C}; \\citealt{march11}) and improve our knowledge of dark energy (e.g., \\citealt{2010MNRAS.406.1759M}; \\citealt{2011MNRAS.418.1707B}). Moreover, the gravitational lensing of SNIa by foreground cosmic structure along their lines-of-sight has been used to constrain cosmological parameters (\\citealt{metcalf99}, \\citealt{dodelson06}, \\citealt{zentner09}) and the properties of the lensing matter (\\citealt{rauch91}, \\citealt{metcSilk}, \\citealt{jacob07}, \\citealt{kronborg10}, \\citealt{jonssonGOODS}, \\citealt{jonssonSNLS}, \\citealt{Karpenka2012}). In addition to their central role in cosmology, the astrophysics of SNIa is also of interest in its own right, and much progress has been made in understanding these objects in recent years (e.g. \\citealt{2000ARA&A..38..191H}).  Other types of SNe are also of cosmological interest.  Type II Plateau Supernovae (SNII-P), for example, can also be used as distance indicators, although only for smaller distances and to lower accuracy than SNIa. Compared to SNIa, however, for which there is still uncertainty regarding the progenitor system, SNII-P explosions are better understood. Furthermore, since SNII-P have only been found in late-type galaxies, biases from environmental effects will most probably have a smaller effect on distance measurements using SNII-P. Thus, the differences between the two types of SNe will result in different systematic effects, allowing SNII-P data to complement SNIa analyses \\citep{2010ApJ...708..661D}.  Although not used directly in cosmology, other classes of SNe are a potential source of contamination when attempting to compile SNIa catalogues, most notably SN Ib/c. The consequences of such contamination have been considered by \\cite{2005ApJ...620...12H}. SN Ib/c are also of considerable astrophysical interest, in particular the nature of their progenitors \\citep{2007PASP..119.1211F}.   The next generation of survey telescopes, such as the Dark Energy Survey (DES; \\citealt{2005ASPC..339..152W}, \\citealt{2011AAS...21743316A}), the Large Synoptic Survey Telescope (LSST; \\citealt{2002SPIE.4836...10T}, \\citealt{2008arXiv0805.2366I}), and SkyMapper (\\citealt{2005AAS...206.1509S}), are expected to observe lightcurves for many thousands of SNe, far surpassing the resources available to confirm the type of each of them spectroscopically. Hence, in order to take advantage of this large amount of SNe data, it is necessary to develop methods that can accurately and automatically classify many SNe based only on their photometric light curves.  In response to this need, many techniques targeted at SNe photometric classification have been developed, mostly based on some form of template fitting (\\citealt{2002PASP..114..833P}; \\citealt{2006AJ....132..756J}; \\citealt{2006AJ....131..960S}; \\citealt{2007AJ....134.1285P}; \\citealt{2007ApJ...659..530K}; \\citealt{2007PhRvD..75j3508K}; \\citealt{2008AJ....135..348S}; \\citealt{2011ApJ...738..162S}; \\citealt{2009ApJ...707.1064R}; \\citealt{2010ApJ...709.1420G}; \\citealt{2010ApJ...723..398F}).  In such methods, the lightcurves in different filters for the SN under consideration are compared with those from SNe whose types are well establised. Usually, composite templates are constructed for each class, using the observed lightcurves of a number of well-studied, high signal-to-noise SNe (see \\citealt{2002PASP..114..803N}), or spectral energy distribution models of SNe. Such methods can produce good results, but the final classification rates are very sensitive to the characteristics of the templates used.  To address this difficulty, \\cite{Newling2011} instead fit a parametrised function  to the SN lightcurves. These post-processed data are then used in either a kernel density estimation method or a `boosting' machine learning algorithm to assign a probability to each classification output, rather than simply assigning a specific SN type.  More recently, \\cite{2012MNRAS.419.1121R} and \\cite{2012arXiv1201.6676I} have introduced methods for SN photometric classification that do not rely on any form of template fitting, but instead employ a mixture of dimensional reduction of the SNe data coupled with a machine learning algorithm. \\cite{2012MNRAS.419.1121R} proposed a method that uses semi-supervised learning on a database of SNe: as a first step they use all of the lightcurves in the database simultaneously to estimate a low-dimensional representation of each SN, and then they employ a set of spectroscopically confirmed examples to build a classification model in this reduced space, which is subsequently used to estimate the type of each unknown SN. Subsequently, \\cite{2012arXiv1201.6676I} proposed the use of Kernel Principal Component Analysis as a tool to find a suitable low-dimensional representation of SNe lightcurves. In constructing this representation, only a spectroscopically confirmed sample of SNe is used. Each unlabeled lightcurve is then projected into this space and a $k$-nearest neighbour algorithm performs the classification.  In this paper, we present a new method for performing SN classification that also does not rely on template fitting in a conventional sense, but combines parametrised functional fitting of the SN lightcurves together with a machine learning algorithm.  Our method is very straightforward, and might reasonably even be described as naive, but nonetheless yields accurate and robust classifications.   The outline of this paper is as follows. In Sec.~\\ref{sec:data}, we describe the data set used for training and testing in our analysis, and we present a detailed account of our methodology in Sec.~\\ref{sec:method}.  We test the performance of our approach in Sec.~\\ref{sec:res} by applying to the data and present our results. Finally, we conclude in Sec.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have presented a new method for performing automated photometric of SNe into Type-Ia and non-Ia. In our a two-stage approach, the SNe lightcurves are first fitted to an analytic parameterised function, and the resulting parameters, together with a few further statistics associated with the fit, are then used as the input feature vector to a classification neural network whose output is the probability that the SN is of Type-Ia. Assuming a canonical threshold output probability $p_{\\rm th}=0.5$, when we train the method using a random sample of 10 (40) per cent of the updated simulated data set released following the SuperNova Photometric Classification Challenge (post-SNPCC), making no selection cuts, we find that it yields robust classification results, namely a completeness of 0.78 (0.82), purity of 0.77 (0.82), and SNPCC figure-of-merit of 0.41 (0.50).  A modest 5--10 per cent improvement in these results is achieved by also including the SN host-galaxy redshift and its uncertainty as inputs to the classification network. The quality of the classification does not depend strongly on the SN redshift.  It is difficult to perform a direct comparison of our results with those submitted to the original SuperNova Photometric Classification Challenge, which are summarised in \\cite{2010PASP..122.1415K}.  As pointed out in that paper, the original challenge data set suffered from a number of bugs; these were subsequently corrected before the release of the post-SNPCC data set used in this paper. The latter also benefited from further improvement in the generation of the simulations, leading to more realistic SNe lightcurves.  It is hard to assess how these differences affect the difficulty of classifying the SNe. For example, in the original SNPCC data set, non-Ia SNe were too dim on average, which made classifcation of Type-Ia SNe easier. Conversely, participants in the original SNPCC were given the spectroscopically confirmed sample $\\mathcal{S}$ of 1,103 SNe and asked to predict the type of SNe in the simulated sample $\\mathcal{P}$ of 20,216 SNe. As discussed in Sec.~\\ref{sec:data}, the fact that $\\mathcal{S}$ is not a representative training sample makes classification more difficult than if one simply uses random training samples, on which most of our analysis has been focused. Despite these caveats, for reference we note that the original challenge entry with the highest SNPCC figure of merit of $\\sim 0.4$ (averaged over SN redshift bins), achieved an overall completeness of 0.96 and purity of 0.79, although the quality of the classification varied considerably with SN redshift.  More recent works by \\cite{Newling2011}, \\cite{2012MNRAS.419.1121R} and \\cite{2012arXiv1201.6676I} analyse the same post-SNPCC data set $\\mathcal{D}$ used in this paper. A meaningful comparison of their results with our own is still not straightforward, however, since all three studies make different choices for the nature and size of the subsets of $\\mathcal{D}$ used for training and testing, which also differ from the choices made in this paper. Nonetheless, some broad comparisons are possible.  The most straightforward comparison is with \\cite{Newling2011}, who present results using, as we do, the subset $\\mathcal{S}$ and also various random subsets of $\\mathcal{D}$ (which they call `representative samples') for training their classifier. For $\\mathcal{S}$, their KDE method achieves a figure-of-merit of 0.37 (0.39) without (with) the inclusion of host-galaxy redshift information, whereas their boosting method yields a figure-of merit of 0.15 using redshift information.  Turning to their analysis of representative samples, from figure~15 in their paper, one sees that their boosting method, which is the more successful of their two methods on the representative samples, achieves figures-of-merit of $\\sim 0.45$ and $\\sim 0.55$ for training sets containing $\\sim 2000$ and $\\sim 8000$ SNe, respectively, which correspond roughly to the sizes of our $\\mathcal{D}_1$ and $\\mathcal{D}_4$ training data sets. These classification results are obtained from the remaining SNe in $\\mathcal{D}$, after including SNe host-galaxy redshift information, and are very similar to the equivalent figures-of-merit achieved by our own classifier, as listed in Table~\\ref{tab:results}. Unfortunately, \\cite{Newling2011} do not give values for their corresponding completeness and purity values, so it is not possible to compare their results with our own.  Both \\cite{2012MNRAS.419.1121R} and \\cite{2012arXiv1201.6676I} adopt very different approaches from the above, and from each other, for choosing the nature and size of the subsets of $\\mathcal{D}$ used for training. Indeed, each of these works considers a range of training sets. \\cite{2012MNRAS.419.1121R} do, however, consider training using the spectroscopically-confirmed sample $\\mathcal{S}$, and obtain a purity of 0.50 (0.54), a completeness of 0.50 (0.90) and a figure-of-merit of 0.13 (0.25) without (with) the inclusion of redshift information. They also construct three further classes of training sets, two of which contain several examples, each requiring the same fixed amount of spectroscopic follow-up time assumed in the construction of the original sample $\\mathcal{S}$. These are: (i) SNe observed in order of decreasing brightness; (ii) ($r$ band) magnitude-limited surveys down to 23.5, 24. 24.5 and 25th magnitude, respectively; and (iii) redshift-limited surveys out to $z=0.4$ and $0.6$, respectively.  In applying their classifier to the post-SNPCC the remaining part of the photometric sample $\\mathcal{P}$, their best classification results are a completeness of 0.65 (0.74), a purity of 0.72 (0.76) and a figure-of-merit of 0.31 (0.36) without (with) the inclusion of redshift information; these were obtained from the deepest magnitude-limited survey, which contained only 165 SNe. It is worth noting the deeper training sets have a SNe class composition that closely resembles that of the full data set $\\mathcal{D}$, and so begin to approximate a representative training sample.  The training sets considered by \\cite{2012arXiv1201.6676I} are closer in spirit to the one used in the original SNPCC. Starting with the spectroscopically-confirmed subsample $\\mathcal{S}$, as a requirement of their method they impose selection cuts such that every SN must have at least one observation epoch with $t \\le t_{\\rm low}$ and one with $t \\ge t_{\\rm up}$ in all available filters. In addition, each SN must have at least 3 observations above a given signal-to-noise ratio (SNR) in each filter. The same selection cuts are also applied to the photometric subsample $\\mathcal{P}$ to produce the corresponding testing sets. Of the selection cuts considered, their sample $D_5$ with $\\mbox{SNR} > 0$ is the least restrictive, yielding training and testing sets containing 830 and 15,988 SNe, respectively. On these demanding data their classifier yields a completeness of 0.44, a purity of 0.37, and a figure-of-merit of 0.06.  From these various comparisons, we conclude that the method presented in this paper is indeed competitive, inspite of its relative simplicity, and yields reasonably robust classifications. Finally, we note that, aside from its relative simplicity and robustness, the classification method that we have presented here can be extended and improved in a number of ways. Firstly, one can easily generalise our approach to divide the non-Ia SNe into their different classes. Moreover, our use of probabilistic classification allows one to calculate expected completeness, purity and figure-of-merit (amongst other measures), without knowing the true classes of the SNe in the testing sample, as will be the case in the classification of real SNe. This allows one to tailor the method by adjusting the output threshold probability $p_{\\rm th}$ to achieve a given completeness, purity or figure-of-merit. Alternatively, one could use the expected ROC curve to \u2018optimise\u2019 the value of $p_{\\rm th}$ used for classification. Nonetheless, one must always remember that if the training sample is not representative of the population, then the predictions for any measure of classification quality will inevitably be biased, in which case the derived threshold probability may, in fact, not be optimal.  We plan to investigate these issues in a future work."
    },
    "1208/1208.3248.txt": {
        "abstract": "Results of observations with the {\\it Spitzer, Hubble, GALEX, Chandra}, and {\\it XMM-Newton} space telescopes are presented for the Luminous Infrared Galaxy (LIRG) merger Markarian 266. The SW (Seyfert 2) and NE (LINER) nuclei reside in galaxies with Hubble types SBb (pec) and S0/a (pec), respectively. Both companions are more luminous than $L^{*}$ galaxies and they are inferred to each contain a $\\approx 2.5\\times10^{8}$ \\Msun\\ black hole. Although the nuclei have an observed hard X-ray flux ratio of $f_X(NE)/f_X(SW)=6.4$, Mrk 266 SW is likely the primary source of a bright Fe K$\\alpha$ line detected from the system, consistent with the reflection-dominated X-ray spectrum of a heavily obscured AGN. Optical knots embedded in an arc with aligned radio continuum radiation, combined with luminous $\\rm H_2$ line emission, provide evidence for a radiative bow shock in an AGN-driven outflow surrounding the NE nucleus. A soft X-ray emission feature modeled as shock-heated plasma with $T \\sim10^7$ K is co-spatial with radio continuum emission between the galaxies. Mid-infrared diagnostics provide mixed results, but overall suggest a composite system with roughly equal contributions of AGN and starburst radiation powering the bolometric luminosity. Approximately 120 star clusters have been detected, with most having estimated ages less than 50 Myr. Detection of 24 \\micron\\ emission aligned with soft X-rays, radio continuum and ionized gas emission extending $\\sim$34\\arcsec\\ (20 kpc) north of the galaxies is interpreted as $\\sim2\\times10^7$ \\Msun\\ of dust entrained in an outflowing superwind. At optical wavelengths this Northern Loop region is resolved into a fragmented morphology indicative of Rayleigh-Taylor instabilities in an expanding shell of ionized gas. Mrk 266 demonstrates that the dust ``blow-out'' phase can begin in a LIRG well before the galaxies fully coalesce during a subsequent ULIRG phase, and rapid gas consumption in luminous dual AGNs with kpc-scale separations early in the merger process may explain the paucity of detected binary QSOs (with pc-scale orbital separations) in spectroscopic surveys. An evolutionary sequence is proposed representing a progression from dual to binary AGNs, accompanied by an increase in observed $\\rm L_{x}/L_{ir}$ ratios by over four orders of magnitude.\\\\ %Force blank line before Keywords; this only works inside the Abstract ",
        "introduction": "\\label{sec:intro}  \\subsection{LIRGs, ULIRGs and GOALS}\\label{subsec:LIRGs}  Luminous Infrared Galaxies (LIRGs; $10^{11}{\\ts}L_\\odot \\le L_{ir} < 10^{12}{\\ts}L_\\odot$) and Ultraluminous Infrared Galaxies (ULIRGs; $L_{ir} \\ge 10^{12}{\\ts}L_\\odot$) are intriguing objects with widespread implications for galaxy evolution. They contain the highest known rates of star formation, they exhibit a high frequency of active galactic nuclei (AGNs) and large-scale outflows (superwinds), and mounting evidence indicates that the majority of local (U)LIRGs represent an evolutionary process involving the transformation of major mergers between dusty, gas-rich disk galaxies into elliptical galaxies hosting classical UV-excess QSOs or powerful radio galaxies \\citep[e.g.,][]{1996ARA&A..34..749S, 2006NewAR..50..701V}. At redshifts of $\\rm z\\sim 1$, LIRGs have a higher space density than ULIRGs and dominate the total star-formation density at that epoch \\citep{2005ApJ...632..169L}. At $\\rm z\\sim2$ the contributions of LIRGs and ULIRGs to the total IR luminosity density are about equal \\citep[e.g.,][]{2007ApJ...660...97C}. Since LIRGs and ULIRGs were much more common in the early universe,  these populations are fundamental in understanding both star formation and galaxy evolution.  The Great Observatories All-Sky LIRG Survey (GOALS)\\footnote{http://goals.ipac.caltech.edu/} is utilizing imaging and spectroscopic data from NASA's {\\it Spitzer, Hubble, Chandra}  and {\\it GALEX} space observatories, in combination with ground-based observations, in a comprehensive study of more than 200 of the most luminous infrared galaxies in the local universe \\citep{2009PASP..121..559A}. The sample consists of 181 LIRGs and 23 ULIRGs that form a statistically complete subset of the flux-limited {\\it IRAS} Revised Bright Galaxy Sample  (RBGS), which itself is comprised of 629 extragalactic objects with 60~\\micron\\ flux densities above 5.24 Jy and Galactic latitudes $\\vert b \\vert > 5^\\circ$ \\citep{2003AJ....126.1607S}.  \\subsection{Dual AGNs in Major Mergers}\\label{subsec:dualAGNs}  A key scientific driver for GOALS is the investigation of the relative time-scales and energetics of active star formation and AGN phenomena during different phases of the merger sequence. It is now widely accepted that all massive galaxies likely have a supermassive black hole (SMBH) in their centers \\citep[e.g.,][]{2005SSRv..116..523F}, with a SMBH mass proportional to the mass of the stellar bulge \\citep[$\\rm M_{SMBH}/M_{bulge}\\approx0.2$\\%;][]{2003ApJ...589L..21M}. However, many uncertainties remain regarding the fueling of paired SMBHs during major mergers. While the nuclei are still separated by kpc scales, how common is it for both SMBHs to have accretion rates high enough to produce energetically significant {\\it dual AGNs}, as opposed to one or both nuclei being powered predominantly by a burst of star formation? How do AGN characteristics depend on properties of the host galaxies and dynamics of the encounter? Are the fuel supplies and accretion rates sufficient to sustain two luminous AGNs well into a true {\\it binary} stage (e.g., binary QSOs), when the SMBHs are closely bound ($\\rm r < 100~pc$) in Keplerian orbits inside a dynamically relaxed (elliptical) merger remnant \\citep[e.g.,][]{2009arXiv0906.4339C}? Is there something special about the physical conditions in systems that host dual AGNs, or is their lack of detection in some (U)LIRGs merely a matter of observing them during the right stage prior to merging (i.e., an evolutionary timing coincidence), or accounting for differences in nuclear dust obscuration? What is the relative importance of AGNs and star formation in the energetics of the ``superwind'' phenomena that appears to be ubiquitous in (U)LIRGs \\citep[e.g., see the review by][]{2006NewAR..50..701V}?  Although (U)LIRGs are predominantly systems involving major mergers, extensive observations utilizing optical and infrared spectroscopy have so far turned up surprisingly few systems containing confirmed dual AGNs. This has a number of possible explanations. First, (U)LIRGs contain large quantities of centrally concentrated dust that can effectively obscure one or both AGNs at optical, near-infrared, and even mid-infrared wavelengths \\citep[e.g.][]{2000A&A...359..887L,2004ApJS..154..184S}. Therefore, circumnuclear star formation can dominate the observed spectra while in reality a powerful AGN may be buried inside. Second,  about 30\\% of observed LIRGs \\citep[e.g.,][]{1995ApJS...98..171V}, and $\\sim$40\\% of ULIRGs \\citep[e.g.,][]{1999ApJ...522..113V} are classified as LINERs based on optical spectroscopy. It has proven difficult to distinguish between low-luminosity AGNs (photoionization from radiation due to accretion onto a SMBH), photoionization from very hot stars, and shock heating (induced by SNe and starburst-driven superwinds) as the primary source of ionization in LINERs \\citep[e.g.,][]{1999ApJ...522..113V}.  Recent results indicate that most nearby LINERs are dominated by photo-ionization rather than shock heating, and that they are an important class of AGNs distinguished primarily by a lower accretion rate than in Seyfert nuclei \\citep{2006MNRAS.372..961K,2008ARA&A..46..475H}. This implies that the frequency of dual AGNs may be much higher than inferred to date. Only recently, with the capability of the {\\it Chandra X-ray Observatory} to penetrate their extensive dust cocoons with high-resolution hard X-ray observations, have suspected dual AGNs been confirmed in a small number of (U)LIRGs. Perhaps the best known example is NGC 6240, in which hard X-rays and strong Fe K$\\alpha$ emission lines indicate the presence of two AGNs  with a projected separation of 1\\farcs4 (1 kpc) \\citep{2003ApJ...582L..15K}. A second case is the ULIRG Mrk 463, a system first pointed out as an active double-nucleus galaxy by \\citet{1978Afz....14...69P}. The eastern component undoubtedly hosts a dust-enshrouded Type 1 AGN that is powering apparently young radio jets \\citep{1991AJ....102.1241M}, but conflicting optical spectral classifications from various studies left the nature of the western nucleus in doubt (LINER, Seyfert 2, or starburst/AGN composite). Dual AGNs in Mrk 463 have been confirmed recently via detection of two compact, luminous sources of hard X-rays using {\\it Chandra} \\citep{2008MNRAS.386..105B}. To date, very few systems in the GOALS sample have been confirmed to contain dual AGNs via X-ray observations. The first system is NGC 6240, and the second is Arp 299 (NGC 3690 + IC 694) \\citep{2004ApJ...600..634B}. This article presents an in-depth investigation of a third system in this rare class of objects that has considerable importance for understanding the evolution of galaxy mergers and their remnants.  \\subsection{Complex Phenomena in Mrk 266}\\label{subsec:background}  One of the most remarkable LIRGs in the local Universe is Mrk 266 (NGC 5256). The system was first called out as an extraordinary object in the Markarian Survey (First Byurakan Sky Survey) of ultraviolet-excess galaxies due to the presence of two nuclei within what appeared to be a single galaxy \\citep{1978Afz....14...69P,1979Afz....15..209P}. In a detailed spectroscopic investigation at optical wavelengths, \\citet{1980Afz....16..621P} showed that the nuclei have a radial velocity difference of 280 $\\rm km~s^{-1}$, and based on the assumption that  the nuclei revolve about a common center of gravity, masses of $\\rm 7\\times10^9~and~10\\times10^9$\\Msun~were estimated for the SW and NE nuclei, respectively. As for all objects in the Markarian Survey, Mrk 266 stood out because its UV-to-optical flux ratio is higher than in normal galaxies. However, Mrk 266 is also a LIRG, with an infrared luminosity of $\\rm L_{ir} = 3.6\\times10^{11} L_\\odot$ \\citep[8 - 1000 \\micron, as estimated from flux densities in all four {\\it IRAS} bands;][]{2009PASP..121..559A}. As shown in the current study (\\S \\ref{subsec:SEDanalysis}), the global spectral energy distribution indicates the bulk of its energy is emitted in the mid- to far-infrared; much more energy is being emitted at the peak of the SED than is escaping in UV radiation ($\\rm \\nu L_{\\nu}(70\\micron)/\\nu L_{\\nu}(0.2\\micron) \\approx 20$).  Mrk 266 contains a filamentary nebula of ionized hydrogen $\\sim$30 kpc in diameter \\citep{1990ApJ...364..471A} surrounding the two nuclei. A luminous X-ray nebula $\\sim$100 kpc in extent detected by {\\it ROSAT} (also surrounding the two nuclei) and complex kinematics derived from optical spectroscopy revealed one of the most spectacular examples of an outflowing, starburst-driven superwind \\citep{1997ApJ...474..659W}. In addition to radio continuum emission from the two nuclei, a bright radio source located between the nuclei was detected and interpreted as enhanced synchrotron emission induced by extensive shocks at the interface of the merging galaxies \\citep{1988ApJ...333..168M}. Imaging in $\\rm H\\alpha + [N II]$ \\citep{1988ApJ...333..168M,1990ApJ...364..471A}, [O III] $\\lambda 5007$ \\citep{1988AJ.....96.1227H} and soft X-rays \\citep{1998AA...336L..21K} also revealed a bright loop of emission extending $\\sim$24\\arcsec (17 kpc) to the north-east and connected to the SW (Seyfert 2) nucleus \\citep{1997ApJ...474..659W}. Recent {\\it Chandra} observations have resolved X-ray emission from both nuclei, and detected diffuse emission associated with the northern feature and between the nuclei \\citep{2007MNRAS.377.1439B}. The nature of the northern emission region (hereafter called the Northern Loop) has remained controversial, with suggestions including a component of the superwind \\citep{1997ApJ...474..659W}, a ``jet'' from an AGN \\citep{1998AA...336L..21K}, and a tidal tail \\citep{2007MNRAS.377.1439B}.  Mrk 266 is worthy of further detailed study because it manifests complex processes that are rarely detected simultaneously, presumably due to their relatively short time scales: vigorous star formation, a candidate dual AGN (depending on the ionizing source of the LINER), a powerful large-scale superwind, and substantial radio continuum and X-ray emission between the nuclei. The relative proximity of the system permits close-up investigation of an important stage in the evolution of major mergers that may involve the transformation of LIRGs into ULIRGs, AGN feedback with the ISM, and implications for fueling (or quenching) of accreting SMBHs which galaxy evolution models predict to have occurred in large numbers during earlier epochs.  In this study, new observations and re-processed archival data from the {\\it Spitzer, Hubble, Chandra, XMM-Newton} and {\\it GALEX} space telescopes are analyzed in combination with multi-wavelength ground-based data. Imaging, photometric, and spectroscopic data are presented in \\S\\ref{sec:Observations}. In \\S\\ref{sec:Discussion} the new data are interpreted to obtain new insights into the nature of the colliding/merging system (\\S\\ref{subsec:Galaxies}), the nuclei and circumnuclear regions (\\S\\ref{subsec:Nuclei}), the region between the nuclei (\\S\\ref{subsec:Center}), the extensive superwind (\\S\\ref{subsec:Superwind}), the SEDs of the major components (\\S\\ref{subsec:SEDanalysis}), newly detected star clusters (\\S\\ref{subsec:SCs}), and the molecular gas properties (\\S\\ref{subsec:MolGas}). In \\S\\ref{subsec:Implications}, Mrk 266 is examined in comparison with other interacting systems with strong radiation sources between the galaxies, and a sequence is proposed that may represent an evolutionary progression from dual AGNs (with kpc-scale separations) to binary AGNs (with parsec-scale orbits). Due to the large number of new measurements spanning many spectral regions, the basic observations and data reductions are described in \\S\\ref{sec:Observations}, but most of the corresponding figures and tables are displayed alongside their interpretation and analyses in \\S\\ref{sec:Discussion}. Conclusions are summarized in \\S\\ref{sec:Summary}. A systemic heliocentric recession velocity of $\\rm 8353~\\pm~13~km~ s^{-1}$ \\citep{1991deVaucouleurs} is adopted, corrected to $\\rm 8825 \\pm~22~km~s^{-1}$ via the flow model of  \\citet{Mould2000a,Mould2000b} that accounts for three major attractors in the local Universe. We adapt cosmological parameters $\\rm H_{\\rm o} = 70~km~s^{-1}~Mpc^{-1}$, $\\Omega_{\\rm M} = 0.28$, and $\\Omega_{\\rm V} = 0.72$ \\citep{2009ApJS..180..225H}. The corresponding luminosity distance to Mrk 266 is 129 Mpc (distance modulus 35.55 mag), and the spatial scale is 0.59 kpc/arcsec\\footnote{\\it Provided by NED at http://ned.ipac.caltech.edu/}.  %%%%%%%% Begin Figure %%%%%%%%%%% %Mrk 266 with HST ACS \\def\\figcapHSTcolor{ \\footnotesize {\\it HST} ACS I band and B band images combined to form a color composite image of the Mrk 266 system and its tidal features. North is up and east is left. The image field of view is 1\\farcm4~x~1\\farcm6, with a scale bar in the lower left corner. The color table has been chosen to maximize contrast for faint extended structures. The limiting surface brightness is 25.0 and 24.5 $\\rm mag~arcsec^{-2}$ in the B and I band, respectively. Labels identify major structures that are studied throughout this article. %\\vskip 0.02truein %Because a page beak does not work to omit a single line of text below this figure } \\ifnum\\Mode=0 %Insert Figure/Table here only in [preprint] or [preprint2] modes \\placefigure{fig:HST_ACS4x} \\begin{verbatim}fig01\\end{verbatim} \\else %For preprint \\ifnum\\Mode=2 \\begin{figure*}[p] \\else \\begin{figure}[p] \\fi \\includegraphics[width=1.0\\textwidth,angle=0]{fig01} \\caption{\\figcapHSTcolor \\label{fig:HST_ACS4x}} \\ifnum\\Mode=2 \\end{figure*} \\else \\end{figure} \\fi %close \\ifnum\\Mode=2 \\fi %close \\ifnum\\Mode=0 %%%%%%%% End Figure %%%%%%%%%%% ",
        "conclusions": "\\label{sec:Summary}  We have presented the results of a multi-wavelength campaign on the luminous infrared galaxy (LIRG) system Markarian 266 that has yielded the following conclusions regarding the morphology, nuclei, dust, molecular gas, outflows, star clusters, and radiation between the nuclei, as well as implications for a better understanding of the origin of ULIRGs and binary AGNs.  1.  {\\it HST} images show that the ``double-nucleus galaxy'' observed in low-resolution images is a strongly interacting pair; we classify the SW and NE galaxies as SBb (pec) and S0/a (pec), respectively. Deeper imaging reveals asymmetric tidal debris spanning $\\approx$100 kpc. At H band, the companions are more luminous than $\\rm L^{*}$ galaxies, and their estimated stellar masses are $\\rm 6.1 \\times10^{10}$ \\Msun\\ (SW) and $\\rm 4.4\\times 10^{10}$ \\Msun\\ (NE). The stellar bulge luminosities imply remarkably similar nuclear black hole masses of $2\\times10^{8}$ \\Msun.  2. The NE (LINER) nucleus is emitting 6.4 times more hard X-ray (2-7 keV) flux than the SW (Seyfert 2) nucleus, yet the luminosity of the NE galaxy is only $\\approx$20\\% of the SW galaxy at 24 \\micron\\ and in CO (1-0) molecular gas emission. The {\\it Chandra} spectra indicate that the SW nucleus is likely the primary source of a bright Fe K$\\alpha$ line detected in the {\\it XMM} spectrum of the total system, consistent with a reflection-dominated X-ray spectrum of a heavily obscured AGN. The F(X-ray)/F([O~III]) flux ratios also indicate the SW nucleus is much more heavily obscured than the NE nucleus.  3. Soft (0.4-2 keV) X-ray emission extends 15 kpc to the north of the system, between the nuclei, and 3 kpc west of the SW nucleus. The smoothed X-ray structure corresponds remarkably well with filaments in H$\\alpha$ and [O III] images. A ridge of X-ray emission between the galaxies has an X-ray spectrum consistent with $\\rm T \\sim 10^7$ K shock-heated gas, strengthening the idea that the corresponding radio emission is shock-induced synchrotron radiation. The lack of detected 9.7 \\micron\\ $\\rm H_2$ S(3) line emission supports the suggestion that a very fast shock at the interface of superwinds expanding from the two galaxies might generate the observed non-thermal radio continuum emission and heat the gas above the dissociation temperature of $\\rm H_2$. The derived cooling time of the X-ray emitting gas between the nuclei is only $\\approx$4 Myr.  4. The HST optical images reveal a circumnuclear arc containing three knots 240 pc west of the NE nucleus. The arc is embedded within B-band emission with a filamentary, bi-conic morphology extending over a radius of 1.2 kpc. This region appears to be dominated by nebular emission from the narrow-line region of a double-sided AGN ionization cone. Radio continuum emission at 1.6 GHz peaks at the optical/near-IR nucleus and has components in alignment with the nucleus and the circumnuclear B-band knots (PA$=$56\\arcdeg). A plausible explanation is that the optical knots are clouds entrained in a shock front produced by an AGN-powered collimated plasma outflow (jet). We liken this region to the radiative bow shock 230 pc south of the LINER in M51.  5. Mid-IR spectral diagnostics suggest the bolometric luminosity is powered by roughly an equal mixture of radiation from AGNs and star formation, with substantial scatter among the various methods. Newly constructed SEDs give infrared luminosities of $\\rm L_{ir} = $ (2.3, 0.7, 3.4)$\\times10^{11}$ \\Lsun\\  for the SW, NE, and total system respectively. Decomposition of the SED of the total system indicates that cold (26 K), cool (72 K), warm (235 K) and hot ($\\sim$1500 K) thermal dust components contribute approximately 57\\%, 20\\%, 12\\%, and 11\\% of the bolometric luminosity, respectively. The total cold dust mass estimate is $\\rm 1.5(\\pm0.4)\\times10^{8}$ \\Msun.  6. {\\it HST} imaging resolves the Northern Loop extending 6-20 kpc from the galaxies into a fragmented morphology suggestive of Rayleigh-Taylor instabilities. A few structures 300-600 pc in length are radially aligned with the NE nucleus, suggestive of bow shocks predicted by hydrodynamic simulations of galactic superwinds. Detection of 24 \\micron\\ emission in the Northern Loop implies that at least $\\rm \\sim2x10^7$\\Msun\\ of dust is being swept out of the system by the superwind during a ``blow-out'' phase that is well underway prior to the final galaxy merger.  7. Approximately 120 star clusters detected with {\\it HST} are concentrated in the SW galaxy and at the base of the Northern Loop; most have estimated ages less than 50 Myr and masses of $\\sim 10^5$ \\Msun. The ratios of cluster surface densities, $\\rm L_{CO}$ and $\\rm L_{ir}$ within 3 kpc of each nucleus are similar (i.e., $\\rm L(SW) / L(NW) \\approx 4-5$). The star cluster luminosity function is similar to what has been found in other LIRGs, and the unobscured clusters contribute little to powering the total infrared luminosity.  8. We conclude that Mrk 266 NE is an AGN based on the presence of: (1) an obscured, hard X-ray point source; (2) a radiative bow shock aligned with ionization cones and apparent radio plasma outflow; (3) PAH emission with small equivalent widths and a deficiency of 6.2 and 7.7 \\micron\\  flux relative to 11.3 \\micron\\  flux; and (4) a $\\rm H_2(1-0)~S(1)~to~Br\\gamma$ flux ratio of 3.5. The spatially extended $\\rm H_2$ emission demonstrates that LINERs can exhibit strong signatures of an AGN and shock excitation.  9. The bulk of the CO (1-0) emission in the system is aligned with the major axis of the SW galaxy. HCN (1-0) is aligned with $\\rm H_2~S(3)$ 9.7 \\micron\\ emission in the nucleus of Mrk 266 NE.  However, the CO (1-0) and HCO$+$ (1-0) emission peaks are located $\\sim$1 kpc SE of the NE nucleus, suggesting that either most of the cold molecular gas has already been stripped from the NE galaxy, or gas is being transferred from the SW galaxy to the NE galaxy. Approximately 40\\% of the total CO (1-0) emission is bridging the galaxies, likely in-falling toward the center of mass or transferring between the galaxies. In this regard, Mrk 266 appears to be in an evolutionary stage similar to VV 114, NGC 6090, and NGC 6240.  10. Two lines of evidence suggest Mrk 266 may evolve into a ULIRG: First, in the context of merger simulations, the galaxies lie within an asymmetric, low surface-brightness halo of tidal debris spanning $\\approx$100 kpc, with a projected separation of 6 kpc and a relative velocity of only $\\rm 135~km~s^{-1}$, indicating they are likely undergoing their second or third encounter with only 50 to 250 Myr remaining until they merge via tidal dissipation. Therefore, Mrk 266 appears to be in a short-lived stage when the nuclei are about to enter the final phase of coalescence characteristic of ULIRGs. Second, the total cold molecular gas mass of $\\rm \\approx7\\times10^{9}$ \\Msun\\ is similar to local and high-redshift ULIRGs. Since $\\approx$40\\% of the total CO (1-0) emission is located between the galaxies, this reservoir is available to form more stars and to fuel the AGNs as the stellar systems and nuclei inexorably coalesce.  11. We propose that Mrk 266 belongs to an evolutionary sequence in which {\\it dual AGNs} with kpc separations represent precursors to putative {\\it binary AGNs} with pc scale orbital radii. In this scenario, where Mrk 266 is in an intermediate phase between Mrk 171 and NGC 6240, the global $\\rm L_x/L_{ir}$ ratio increases by over four orders of magnitude as obscuring material is expelled by outflows to gradually expose previously obscured AGNs.  12. Since major mergers provide a natural process to form SMBH/AGN pairs, the scarcity of confirmed and candidate binary QSOs in large spectroscopic and VLBI surveys is unexpected. While two possibilities are raised in the literature, that either the SMBHs inspiral very rapidly, or fueling of the accretion disks is quenched during the binary phase, we propose a third hypothesis. Since the gas-depletion lifetime of ULIRGs and LIRGs is 10-100 times shorter than the time-scale for creation of a binary SMBH ($\\ga$1 Gyr), it is possible that, in most instances, the gas will be consumed by star formation and accretion during a dual AGN phase long before the SMBHs inspiral to sub-pc separation.  %\\ifnum\\Mode=2 %\\pagebreak %\\fi"
    },
    "1208/1208.6502_arXiv.txt": {
        "abstract": "The relative age of galaxies at different redshifts can be used to infer the Hubble parameter and put constraints on cosmological models. We select luminous red galaxies (LRGs) from the SDSS DR7 and then cross-match it with the MPA/JHU catalogue of galaxies to obtain a large sample of quiescent LRGs at redshift $z\\sim 0.03-0.39$.  The total 23,883 quiescent LRGs are divided into four sub-samples according to their velocity dispersions and each sub-sample is further divided into 12 redshift bins. The spectra of the LRGs in each redshift and velocity bin are co-added in order to obtain a combined spectrum with relatively high $S/N$.  Adopting the GalexEV/SteLib model, we estimate the mean ages of the LRGs from these combined spectra by the full-spectrum fitting method. We check the reliability of the estimated age by using Monte-Carlo simulations and find that the estimates are robust and reliable. Assuming that the LRGs in each sub-sample and each redshift bin were on average formed at the same time, the Hubble parameter at the present time $H_0$ is estimated from the age--redshift relation obtained for each sub-sample, which is compatible with the $H_0$ value measured by other methods. We demonstrate that a systematic bias (up to $\\sim 20\\%$) may be introduced to the $H_0$ estimation because of recent star formation in the LRGs due to the later major mergers at $z\\la 0.4$, but this bias may be negligible for those sub-samples with large velocity dispersions. Using the age--redshift relations obtained from the sub-sample with the largest velocity dispersion or the two sub-samples with high velocity dispersions, we find $H_0= 65^{+7}_{-3}\\kmsmpc$ or $H_0= 74^{+5}_{-4}\\kmsmpc$ by assuming a spatially flat $\\Lambda$CDM cosmology. With upcoming surveys, such as the Baryon Oscillation Spectroscopic Survey (BOSS),  even larger samples of quiescent massive LRGs may be obtained, and thus the Hubble parameter can be measured with high accuracy through the age--redshift relation. ",
        "introduction": "The expansion history of the universe are presently studied with a few observational probes, such as  the supernova Ia, baryon acoustic oscillations (BAO), weak gravitational lensing, and galaxy clusters, etc.  Each of these probes has its pros and cons, and suffer from different systematic uncertainties \\citep[e.g.,][]{Fre10}.  A new observational probe of the cosmic expansion history would be invaluable, and can provide additional cross check with the results obtained from the existing methods. Combining the results obtained by different means may further help to constrain robustly the dynamical nature of the universe.  \\citet{Jim02} proposed a novel approach to explore the expansion history of the universe, which is based on the age--redshift relation of passively evolving massive galaxies. Assuming that the passively evolving galaxies at different redshifts were born approximately at the same time, the age of these galaxies can then be taken as a cosmic chronometer. If the ages of such galaxies can be accurately estimated, then this age--redshift relation may be used to determine the cosmic expansion history.  Even if there is some systematic errors in the absolute age measurements, it is argued that such errors could be canceled in the relative age of these galaxies at different redshifts, thus providing a good measurement of $H(z)$: \\begin{equation} H(z) = -\\frac{1}{1+z}\\frac{dz}{dt}. \\label{eq:dtdz} \\end{equation} Indeed, observations show that the most massive galaxies are mainly composed of old stellar populations formed at redshifts $z > 1-2$, less than $1\\%$ of their present stellar mass is formed at $z<1$ \\citep{Dun96,Spi97,Cow99, Hea04,Tho05,Cim08,Tho10}, hence these galaxies are suitable for this application.  \\citet{Jim03} applied this method to a sample of massive galaxies at low redshift by fitting their spectra with the single stellar population (SSP) spectra based on the SPEED model \\citep{Jim04}, and obtained $H_0=69 \\pm~12\\kmsmpc$.  \\citet{Sim05} assembled a high redshift data set obtained from the Gemini Deep Survey (GDDS) and some other archival data, and applied the same method to estimate $H(z)$ for a large redshift range ($z\\sim 0.1-1.8$). These earlier works on the age--redshift relation adopted the SSP to fit each galaxy spectrum in the sample, and selected the age of the oldest one in each redshift bin as the envelop of the age. However, the poor signal-to-noise ratio ($S/N$) spectra of individual galaxies and the contamination from the telluric emission and absorption may lead to uncertainties in the age estimates, and the method of the oldest galaxy envelop draw results from a small number of galaxies at the extremes of the distribution, which may also undermine the validity of the result, and makes the method hard to use.  To overcome this problem, \\citet{Car10} obtained the combined spectra for those luminous red galaxies (LRGs) with similar physical properties in each redshift bin by co-adding their spectra, of which the $S/N$ is much higher than individual galaxies. They then estimated the age of the combined spectra by using the standard Lick absorption line indices, which may be regarded as the mean age of a large sample of galaxies. They obtained the age--redshift relation, but they did not use this relation to further constrain the Hubble parameter.  In this paper, we first select a LRG sample from the SDSS data release 7 (DR7). In order to improve the $S/N$ and remove the contamination, we also use the combined spectrum rather than the single spectrum of each galaxy.  However, we adopt the full spectrum fitting method, different from the standard Lick absorption line indices adopted by \\citet{Car10}, to estimate the mean age of the combined spectrum, and then obtain the age--redshift relation. Furthermore, we also use the age-redshift relation obtained from the combined spectra to constrain the Hubble parameter at the present time $H_0$ and analyze the possible systematic bias in the estimated $H_0$.  The paper is organized as follows. In Section~2, we describe the selection criteria of the LRG sample. In Section~3, we provide the details of the fitting method and the age--redshift relation estimated from the LRG sample. In Section~4, we constrain the Hubble parameter by using the obtained age-redshift relation. Discussions on the resulted age--redshift relation and the possible associated systematic bias are given in Section~5.  Conclusions are summarized in Section~6. ",
        "conclusions": "\\label{sect:conclusion}  In this paper, we selected $23,883$ quiescent LRGs from the SDSS DR7 in the redshift range from $0.03$ to $0.39$, by setting a threshold of zero emission (at a $2-\\sigma$-level) of the H$_\\alpha$ and [O{\\sc{ii}}] lines directly obtained from the MPA/JHU catalogue. The quiescent LRG sample is divided into four sub-samples according to galaxy velocity dispersions. For each sub-sample, the spectra of galaxies in each of the $12$ redshift bins (from $z=0.03$ to $0.39$ with a step of $\\delta z=0.03$) are combined together to obtain a high $S/N$ combined spectrum. Using the full spectrum fitting method, the luminosity-weighted physical properties, such as the velocity dispersion, the metallicity and the age, of those quiescent LRGs are obtained from the combined spectra by adopting a single population synthesis model, i.e., the GalexEV/SteLib model. Using Monte-Carlo simulations, we find that the model results are robust and reliable. We argue that the age--redshift relation estimated from the LRG sample could be systematic biased because of the contamination from a possible younger stellar population formed at $z\\la 0.4$ as consequence of major mergers. This bias is most significant for LRGs with smaller velocity dispersions but insignificant for the most massive LRGs. Considering of this systematic bias, the age--redshift relation obtained from the model fittings is fully consistent with the expectations from the $\\Lambda$CDM cosmology.  The Hubble parameter $H_0$ is first estimated by using the age--redshift relation obtained from each sub-sample, and its value ranges from $89^{+7}_{-9} \\kmsmpc$, $83^{+9}_{-8}\\kmsmpc$, $72^{+6}_{-7}\\kmsmpc$, to $65^{+7}_{-3}\\kmsmpc$ for the four sub-samples with velocity dispersions from low to high, respectively. The large value of the $H_0$ estimated from the sub-samples with low velocity dispersion is probably due to the systematic bias, which can be as high as $\\sim 20\\%$. Using the age--redshift relations obtained from the two sub-samples with high velocity dispersions or the sub-sample with the largest velocity dispersion, we find $H_0= 74^{+5}_{-4}\\kmsmpc$ or $H_0= 65^{+7}_{-3}\\kmsmpc$ if assuming a spatially flat $\\Lambda$CDM cosmology, which may be less affected by the systematic bias, close to the true $H_0$, and are well consistent with the best estimates through other techniques. However, it needs further test on whether those most massive galaxies are truly passively evolving or not.   In summary, we have demonstrated that the age--redshift relation of quiescent galaxies can be reliably estimated by using the full spectral fitting method if the $S/N$ of their spectra are sufficiently high. We conclude that some cosmological parameters, such as the Hubble parameter, can be constrained with considerable accuracy through the age--redshift relation obtained from those most massive LRGs, which is totally independent of other methods. With future surveys like BOSS, the Hubble parameter may be tightly constrained by the age--redshift relation obtained from the most massive quiescent LRGs."
    },
    "1208/1208.4808_arXiv.txt": {
        "abstract": "We present a deep, wide-field optical study of the M81 group dwarf galaxy Holmberg~II (HoII) based on Subaru/Suprime-Cam imaging. Individual stars are resolved down to $I\\sim25.2$, i.e. about 1.5~mag below the tip of the red giant branch (RGB). We use resolved star counts in the outskirts of the galaxy to measure the radial surface brightness profile down to $\\mu_V\\sim 32$ mag arcsec$^{-2}$, from which we determine a projected exponential scalelength of $0.70\\arcmin\\pm0.01\\arcmin$ (i.e. $0.69\\pm 0.01$~kpc). The composite profile, ranging from the cored centre out to R=7$\\arcmin$, is best fit by an EFF profile which gives a half-light radius of $1.41\\arcmin\\pm0.04\\arcmin$ (i.e. $1.39\\pm 0.04$~kpc), and an absolute magnitude M$_V=-$16.3. The low surface-brightness stellar component of HoII is regular and symmetric and has an extent much smaller than the vast \\hi\\ cloud in which it is embedded. We compare the spatial distribution of the young, intermediate age, and old stellar populations, and find that the old RGB stars are significantly more centrally concentrated than the young stellar populations, contrary to what is observed in most dwarf galaxies of the Local Universe. We discuss these properties in the context of the comet-like distribution of \\hi\\ gas around HoII, and argue for the presence of a hot intragroup medium in the vicinity of HoII to explain the contrasting morphologies of the gas and stars. ",
        "introduction": "\\label{sec:1}  In spite of their intrinsic faintness and minimal contribution to the total light, the stellar outskirts of galaxies hold crucial information about the processes of galaxy formation and evolution. Interactions, mergers and accretions all leave their imprint on the outer stellar populations in the form of substructures, streams and diffuse haloes. In addition, the long evolutionary timescales and high sensitivity to external influences means that coherent substructures are easier to detect and longer-lived than in the denser inner regions.  Considerable effort has been devoted to studying the outer regions of massive galaxies \\citep[see, e.g.][and references therein]{bar12} which, according to hierarchical models of galaxy formation, have acquired a significant fraction of their mass through mergers and accretion episodes.  The stellar peripheries of low mass dwarf galaxies have been much less studied. Most work to date has focused on Local Group (LG) galaxies (e.g. Fornax: \\citealt{col05}; NGC6822: \\citealt{deb06}; Sculptor: \\citealt{wes06}), although a few dwarfs in nearby groups have also been targeted \\citep[e.g.][]{rys11}.  The strong morphology-density relation exhibited by dwarf galaxies suggests external mechanisms play a major role in shaping their evolution \\citep[e.g.][]{wei11}; surveying their stellar outskirts may therefore yield clues on the dominant processes involved.  The global picture emerging from previous work is that most dwarf galaxies have a smooth, and generally old and metal-poor `halo' surrounding a more concentrated, younger and/or more metal-rich population.  While the stellar haloes of large galaxies are believed to have formed from the accretion of smaller galaxies at high redshift, it is unclear whether this process has been significant in dwarfs.  For example, pure accretion cannot explain the main properties of dwarf galaxies' haloes, i.e., smooth distribution, and the existence of age and/or metallicity gradients. Instead, these observations suggest either a `shrinking' scenario, in which the region of active star formation contracts as gas supply diminishes \\citep{hid09,zha12}, or radial migration of stars formed close to the centre towards the outskirts \\citep{sti09}.  Here we analyse the stellar outskirts of the dwarf galaxy Holmberg\\,II (hereafter abbreviated as HoII).  HoII is a dwarf irregular galaxy in the M81 group that was discovered by \\citet{hol50} while surveying the galaxies in this group. Due to its location on the near-side of the group ((m$-$M)$_0$=27.65, i.e. 3.4~Mpc), as well as its proximity to the Sc spiral galaxy NGC2403 and similar radial velocity, it is usually associated with the NGC2403 subgroup along with three other dwarf irregular galaxies \\citep{kar02}. HoII is very similar to the Small Magellanic Cloud (SMC) in terms of absolute magnitude, \\hi\\ and total mass: M$_B\\sim-$16.7, M$_{\\rm HI} \\sim 6 \\times 10^8 M_{\\sun}$, and M$_{\\rm tot} = 2.1 \\times 10^9 M_{\\sun}$ \\citep{wal07,oh11} compared to M$_B\\sim-$16.1, M$_{\\rm HI} \\sim 4 \\times 10^8 M_{\\sun}$, and M$_{\\rm tot} = 2.4 \\times 10^9 M_{\\sun}$ for the SMC \\citep{sta99,sta04}.  \\defcitealias{bur02}{BC02}  \\begin{figure} \\includegraphics[width=8.3cm]{f1.eps} \\caption{Mosaic V image of HoII, showing the whole field-of-view of our Subaru data ($\\sim$35.8$\\arcmin\\times$29.3$\\arcmin$). The \\hi\\ contours from \\citetalias{bur02}, ranging from $N_{HI}= 0.1$ to $19\\times 10^{20}$ atoms cm$^{-2}$, are overplotted. Note the comet-like morphology of the outermost contour.\\label{fig1}} \\end{figure}  While HoII has been observed at all wavelengths, to date the only deep resolved stellar populations study comes from HST/ACS observations \\citep{wei09} which cover a relatively small fraction of the galaxy. The need for deep wider data is especially motivated by the striking morphology of its \\hi\\ cloud. From deep VLA data, \\citet[hereafter BC02]{bur02} found that the distribution of neutral hydrogen has a cometary appearance -- compressed on one side with a faint extended component on the opposite side -- with the tail pointing away from the centre of the M81 group. \\citetalias{bur02} argued that HoII is moving toward the M81 group and that ram pressure from a hot intragroup medium (IGM) is responsible for the \\hi\\ morphology, although they could not rule out the alternative interpretation of a gravitational interaction between HoII and one of its fainter neighbours.  \\begin{figure*} \\includegraphics[width=15.0cm]{f2.eps} \\caption{Color composite mosaic of HoII from our Subaru data. The image is cropped to $\\sim$14$\\arcmin$ on a side; North is up and East to the left. \\label{fig2}} \\end{figure*}  Possible signatures of ram pressure stripping have been observed in a number of galaxies of the local universe \\citep[e.g.][]{con83,gav95,ryd97,ken04,chu07,mcc07}. However, most of these galaxies are too distant to be resolved into individual stars and their analyses have been confined to fairly high surface brightness inner regions. Knowledge of the distribution of resolved stellar populations at very large radii is fundamental because the stars do not respond to ram pressure. On the other hand, tidal forces affect gas and stars equally so both components should exhibit similar asymmetries.  Comparing the large-scale distribution of the stars with that of the gas in HoII has the potential to reveal whether the \\hi\\ morphology was caused by ram pressure, tidal forces, or a combination of both.  In this paper we present a deep, wide-field study of HoII based on Subaru/SuprimeCam data, and analyse the properties of its stellar populations in the context of the comet-like shape of the \\hi\\ cloud in which it is embedded.  In Section~\\ref{sec:2}, we describe the observations and data reduction, and present the resulting CMDs in Section~\\ref{sec:3}. The spatial distribution of the various stellar populations are decribed in Section~\\ref{sec:4}. In Section~\\ref{sec:5}, we present the radial profile and constrain the spatial extent of HoII. We discuss the implications of our results regarding the peculiar stellar and \\hi\\ distributions in Section~\\ref{sec:6}, and summarize the main results in Section~\\ref{sec:7}. ",
        "conclusions": "\\label{sec:7}  We have carried out a wide-field survey of the M81 group dwarf galaxy Holmberg~II based on deep Subaru/Suprime-Cam imaging in $V$ and $I$. These observations cover the whole extent of the galaxy, including the vast \\hi\\ cloud, and allow us to perform photometry of individual stars down to $I\\sim25.2$, i.e. about 1.5~mag below the tip of the RGB.  The deep CMDs reveal the presence of stellar populations of all ages, from a few Myr old (MS+BSG, RSG) to several Gyr old (RGB).  While in most dwarf galaxies in the Local Universe the younger stars are found to be more centrally concentrated than the older populations \\citep[e.g.][]{zha12}, we find that in HoII the old RGB stars are significantly more concentrated than the young MS+BSG stars. Indeed, we find that the exponential scalelength for the young MS+BSG population is much larger than that of the RGB component (2.8$\\arcmin\\pm0.4\\arcmin$ vs.\\ 0.76$\\arcmin\\pm0.04\\arcmin$, respectively). We speculate that the shockwave due to ram pressure increased the gas density in the central part of the \\hi\\ cloud and triggered star formation on large scales.  Our Subaru data enable us to construct a composite surface brightness profile for HoII by combining diffuse light in the central region with star counts at large radii.  This profile is one of the deepest yet published for any galaxy, extending from the centre out to ${\\rm R} \\sim$~7$\\arcmin$ where $\\mu_V=32$ mag arcsec$^{-2}$.  Fitting an exponential profile to the outer regions gives a (projected) scalelength of $0.70\\arcmin\\pm0.01\\arcmin$, corresponding to $0.69\\pm 0.01$~kpc at the distance of HoII.  Finally, we discuss the properties of the resolved stellar populations in the context of the morphology of the large \\hi\\ cloud in order to understand the origin of its swept-back, cometary appearance. Previous studies based on diffuse optical light or 21~cm data could not definitively determine whether the cloud shape was due to ram pressure from a hot IGM or to a tidal interaction with a nearby companion galaxy.  Our deep photometry shows that the intermediate-age and old stars have a regular circular distribution and show no sign of tidal tails/streams. In addition, we find that there are very few, if any, HoII stars beyond R$\\sim$~7$\\arcmin$ where the \\hi\\ becomes distorted. Since tidal forces would affect gas and stars equally, our data strongly suggest that the spectacular morphology of the \\hi\\ cloud is due to ram pressure. The detection of significant amount of diffuse hot gas in the vicinity of HoII would further verify this."
    },
    "1208/1208.4663_arXiv.txt": {
        "abstract": "We present high-resolution spectroscopy of gaseous CO absorption in the fundamental ro-vibrational band toward the heavily obscured active galactic nucleus (AGN) IRAS~08572$+$3915. We have detected absorption lines up to highly excited rotational levels ($J\\leqq$ 17). The velocity profiles reveal three distinct components, the strongest and broadest ($\\Delta v$ $>$ 200 km~s$^{-1}$) of which is due to blueshifted ($-$160 km~s$^{-1}$) gas at a temperature of $\\sim$ 270 K absorbing at velocities as high as $-$400 km~s$^{-1}$. A much weaker but even warmer ($\\sim$ 700 K) component, which is highly redshifted ($+$100 km~s$^{-1}$), is also detected, in addition to a cold ($\\sim$ 20 K) component centered at the systemic velocity of the galaxy. On the assumption of local thermodynamic equilibrium, the column density of CO in the 270 K component is $N_{\\rm{CO}}\\sim4.5\\times10^{18}$ $\\rm{cm^{-2}}$, which in fully molecular gas corresponds to a $\\rm{H_2}$ column density of $N_{\\rm{H_2}}\\sim2.5\\times10^{22}$ $\\rm{cm^{-2}}$. The thermal excitation of CO up to the observed high rotational levels requires a density greater than $n_c(\\rm{H_2})>2\\times10^{7}$ $\\rm{cm^{-3}}$, implying that the thickness of the warm absorbing layer is extremely small ($\\Delta d<4\\times10^{-2}$ pc) even if it is highly clumped. The large column densities and high radial velocities associated with these warm components, as well as their temperatures, indicate that they originate in molecular clouds near the central engine of the AGN. ",
        "introduction": "\\label{sec:intro}  Active galactic nuclei (AGNs) are broadly classified into two types: type 1 AGNs, which display broad permitted emission lines in the optical spectrum; and type 2 AGNs, which display both permitted and forbidden lines with narrow line widths. The detection of broad emission lines in polarized light from type 2 AGNs (Antonucci \\& Miller 1985) present a clue to tie these two types together. The AGN activity is powered by a supermassive black hole ($10^6$--$10^9$ $\\MO$) and its accretion disk, which extends to $\\leq$1 pc. This central engine is surrounded by an optically and geometrically thick dusty torus, extending to $\\sim$ 1--100 pc. In the ``unified model'' of AGNs (e.g., Antonucci 1993), it is postulated that much of the variety in types of AGNs is not due to intrinsic differences between them, but is largely the result of their strongly non-spherical geometries and the different orientations of their central regions with respect to our lines of sight. Specifically, in a type 1 AGN the torus is observed face-on and the interior may be observed directly, while in a type 2 AGN the torus is observed edge-on, which prevents a direct view of the nucleus. The obscuring torus around the central engine is the key element for the AGN unified model, and the strongest verification of it would come from the detection of the molecular torus itself. Recent observations at many wavelengths, for example the detection of silicate emission from type 1 AGNs (e.g., Hao et al., 2007), surely show the presence of obscuring tori surrounding the nuclei. However, the physical conditions and structure of the torus have never been measured directly, and we have only a crude idea of them. The direct measurements of the physical conditions (temperature and column density) of the molecular torus and the determination of the gas kinematics within the torus are long awaited goals of extragalactic astronomy.  Carbon monoxide (CO) is the most abundant interstellar molecule after hydrogen molecule ($\\rm{H_2}$). For the study of molecular clouds, pure rotational emission lines of CO have been extensively observed in the millimeter and sub-millimeter wavelength regions. However, to date such observations have not been suitable for the detailed characterization of distant molecular clouds around the AGNs, because they have provided information on neither the detailed spatial structure (due to the relatively large beam sizes of current millimeter telescopes) nor the physical conditions (due to the limited number of observable lines) of the clouds. In addition, millimeter observations can be contaminated by large-scale CO emission in the host galaxies on scales of 10--100 kpc, which can make it difficult to clearly distinguish the molecular gas in close proximity to the AGN, i.e. the putative molecular torus.  The technique we employ in the current work is high-resolution spectroscopy of absorption lines of the CO fundamental ($v$=1$\\leftarrow$0) ro-vibration band, whose center lies near 4.7 $\\rm{\\mu m}$. Continuum emission associated with the bright, compact central region of an AGN is used as a background source and the foreground molecular clouds are to be observed in absorption. This technique achieves very high spatial resolution, since the resolution is determined by the size of the background continuum source. In addition, a large number of CO lines covering a wide range of rotational levels can be observed simultaneously with the same instrument under the same conditions. This enables accurate determination of the physical conditions, such as temperatures and column densities of the absorbing molecular clouds.  In spite of these advantages, most previous ground-based observations of the 4.7 $\\rm{\\mu m}$ CO absorption spectrum have been limited to Galactic sources (e.g., Scoville et al. 1983; Mitchell \\& Maillard 1993; Nakagawa et al. 1997; Goto et al. 2003), with only a few examples toward extragalactic objects (Spoon et al. 2003; Geballe et al. 2006). This is mainly due to difficulties in achieving high sensitivities at the wavelength around 4.7 $\\rm{\\mu m}$. 8-m class telescopes with high sensitivity, now are making it possible to observe the CO ro-vibrational absorption lines in molecular clouds around some AGNs.  One class of AGNs whose central nuclei are obscured by gas and dust in our line of sight, the so-called Seyfert 2 galaxies, would appear to be the most suitable for such observations. However, spectra of objects in this class have to date shown no significant absorption by the CO fundamental ro-vibration band (Lutz et al. 2004), in some cases down to very low limits (Mason et al. 2006; Geballe et al. 2009). On the other hand, the {\\it Spitzer Space Telescope} has detected a strong absorption by gaseous CO toward the dusty Ultra-Luminous InfraRed Galaxy (ULIRG), IRAS~F00183$-$7111 (Spoon et al. 2004). Observations at the VLT by Risaliti et al. (2006) and Sani et al. (2008) also have detected CO fundamental band absorption toward some ULIRGs whose central sources are heavily obscured. These results suggest that CO absorption does not appear in typical Seyfert 2 galaxies but can be found in at least some heavily obscured AGNs. The reasons for this are unclear, although in the case of the prototypical Seyfert 2 galaxy, NGC~1068, Geballe et al. (2009) have concluded that most of the gas in front of its AGN is diffuse in nature and therefore contains little CO. Both classes of objects, Seyfert 2 galaxies and ULIRGs, appear to have considerable obscuring material near their AGNs. However, hard X-ray observations (Franceschini et al. 2003; Ptak et al. 2003) suggested that AGNs in ULIRGs are not necessarily the dominant source of nuclear luminosity and that nuclear star formation is a significant contributor. Ongoing star formation in the nuclei of ULIRGs implies the presence of large amounts of molecular gas there.  ULIRGs (Sanders \\& Mirabel 1996) are galaxies whose extremely large, quasar-like luminosities ($L_{\\rm{IR}} \\geq 10^{12} \\LO$) emerge at least 90 percent in the form of infrared dust emission. They are associated with interacting or merging gas-rich galaxies (Armus et al. 1987; Sanders et al. 1988; Leech et al. 1994; Clements et al. 1996; Murphy et al. 1996), in which the collision produces either an extreme burst of star formation (usually near the nucleus of one of the merging galaxies), a greatly increased rate of infall into the vicinity of an AGN, or both. In many cases it is difficult to distinguish between the two phenomena, because the large amount of dust prevents direct observation of the energy source. However, ULIRGs with bright pointlike nuclei at mid-infrared wavelengths, whose spectra show strong dust absorption, are generally thought to possess luminous buried AGNs (Imanishi, Dudley \\& Maloney 2006). In the case of a ULIRG-producing merger involving a galaxy that contains an AGN, the survival of the molecular torus in the AGN is unclear. The torus could be disrupted by gas falling into the nucleus or by ejection of gas caused by the rapid increase in luminosity of the AGN precipitated by such infall, or simply by intense heating by the AGN alone.  Infrared spectroscopy of ULIRGs is a possible way to investigate the energetic events in the central region of obscured AGNs. For such purpose, the observations of the CO fundamental band absorption toward ULIRGs with space telescopes, Spitzer and AKARI, are effective in providing some information (Spoon et al. 2005, Imanishi et al. 2008, Imanishi et al. 2010). However, the spectral resolutions of those facilities are relatively low ($R$ $\\sim$ 600 for the Spitzer/IRS, $R$ $\\sim$ 100 for the AKARI/IRC) and are insufficient to resolve molecular ro-vibrational bands into individual transitions, let alone resolve velocity structure. Thus, in order to reveal the physical conditions of the circumnuclear region and the origin of the CO absorption, we have obtained high-resolution spectra toward heavily obscured AGNs, using ground-based telescopes.  In this paper we report on the high resolution 4.90--5.13 $\\rm{\\mu m}$ spectrum toward the obscured AGN IRAS~08572$+$3915 NW, using the 8.2-m Subaru Telescope. The detection of CO absorption lines in the fundamental band toward this galaxy was reported by Geballe et al. (2006), using the United Kingdom 3.8-m Infrared Telescope (UKIRT). Although the signal-to-noise ratio of their spectrum is modest, they detected broad and complex absorption line profiles. Their spectrum shows that the line profile are dominated by a very broad absorption feature, which extends from 0 to $-$400 km~s$^{-1}$ and peaking near $-$160 km~s$^{-1}$ relative to the systemic velocity. They also detected a second velocity component close to the systemic velocity, and marginally detected a third component that is redshifted. Our new and improved spectrum covers a greater portion of the band and has a considerably higher signal-to-noise ratio. Using it we can clearly separate the CO absorption into the three components deduced from the earlier spectrum. The new spectrum can be used to better constrain the physical conditions in the absorbing molecular clouds.  The characteristics of IRAS~08572$+$3915 are reviewed in the next section. The observations and the data reduction procedures are described in Section~\\ref{sec:obs}. In Section~\\ref{sec:results}, we present the observed spectrum and the line profiles. Temperatures and column densities of the absorbing gas are estimated in Section~\\ref{sec:tmpcol} assuming local thermodynamic equilibrium (LTE). The origin of each velocity component of the absorbing gas is discussed in Section~\\ref{sec:component}, and the line of sight dimensions of the warm molecular clouds is estimated in Section~\\ref{sec:size}. The locations of the continuum emitting and absorbing regions are discussed in Section~\\ref{sec:cont}. In Section~\\ref{sec:discussion}, we discuss the relation between the warm molecular clouds we discovered and the putative AGN molecular torus. A summary is given in Section~\\ref{sec:summary}. ",
        "conclusions": "\\begin{tabular}{lcccc} \\hline Covering factor & Component & Temperature & CO column density & $\\rm{H_2}$ column density \\\\ &  & (K) & ($10^{18}$ $\\rm{cm^{-2}}$) & ($10^{22}$ $\\rm{cm^{-2}}$) \\\\ \\hline  $C_f=1.0$...............& Cold & $23\\pm1$ & $0.574\\pm0.006$ & $0.319\\pm0.003$ \\\\ $C_f=1.0$...............& Warm & $325\\pm5$ & $2.65\\pm0.04$ & $1.47\\pm0.02$ \\\\ $C_f=0.6$...............& Cold & $24\\pm1$ & $1.98\\pm0.01$ & $1.10\\pm0.01$ \\\\ $C_f=0.6$...............& Warm & $273\\pm2$ & $4.48\\pm0.04$ & $2.49\\pm0.02$ \\\\ \\hline \\\\  \\multicolumn{5}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize Notes. \\par\\noindent Col.(4): Total CO column density derived on the basis of the temperature in Col.(3). \\par\\noindent Col.(5): Total $\\rm{H_2}$ column density, assuming the standard cosmic abundance ratio of $N_{\\rm{CO}}/N_{\\rm{H_2}}\\sim1.8\\times10^{-4}$ (Dickman 1978). }\\hss}}  \\end{tabular} \\end{center} \\end{table*}  We caution that because the CO absorbs over an extremely wide velocity range the assumption that $C_f$ is independent of velocity, based only on analysis of the low-$J$ lines, whose equivalent widths have comparatively large contributions from the cold gas, is questionable. It also is not obvious that $C_f=0.6$ can be applied to all the lines. Thus there is some uncertainty in the estimates of the column density. Note, however, that the difference between the results with $C_f=0.6$ and 1.0 is not large for the warm component, which is the main topic of the present paper. Hence, we take $C_f=0.6$ as the nominal value for the remainder of the paper.    \\subsection{Origin of each velocity component}\\label{sec:component}  Our observations have clearly shown the existence in front of the northwestern nucleus of IRAS~08572$+$3915 of three basic gaseous components associated with different velocity ranges centered at the systemic velocity and centered at $-$160 km~s$^{-1}$ and $+$100 km~s$^{-1}$ relative to it, and with three widely different excitation temperatures, only two of which are apparent in Figure~\\ref{fig:BoltzmanC6}. In this subsection we discuss the origin of each velocity component. We have not been successful at cleanly isolating these components in velocity space for detailed analysis and comparison, using standard deconvolution techniques, because the velocity profiles are too complicated. The velocity extents, velocity ranges, and degrees of overlap of the components vary with rotational level. Moreover, the profile of each velocity component is asymmetric, and each could consist of several sub-components. Nevertheless, the large temperature difference of the two dominant absorption components allows us to use Figure~\\ref{fig:LineShape} and Figure~\\ref{fig:BoltzmanC6} to crudely separate and simply characterize them. The much weaker third component differs sufficiently in temperature and velocity from the others that it can also be crudely characterized.  The cold component at $\\sim$24 K revealed by Figure~\\ref{fig:BoltzmanC6} must correspond to CO absorption at velocities near the systemic velocity. As seen in Figure~\\ref{fig:LineShape}, this component was detected only in transitions from levels with $J\\leqq5$, implying a gas temperature around 30 K. We interpret this component as arising in interstellar clouds in the host galaxy located far from the nucleus, since its excitation temperature is low and is similar to those of giant molecular clouds in the Galaxy. The presence of a strong 3.4 $\\mu$m carbonaceous dust absorption feature toward IRAS~08572$+$3915 (e.g., Imanishi, Dudley \\& Maloney 2006) suggests that a considerable fraction of the interstellar molecular gas along the line of sight is in diffuse clouds, as is the case toward the Galactic center (Whittet et al. 1997). In diffuse clouds almost all of the carbon ($\\sim$ 99\\%) is atomic and there very little of it is in CO ($\\sim$ 1\\%). Thus, if the gas traced by the observed cold CO component is entirely in diffuse clouds, the $\\rm{H_2}$ column density could be two orders of magnitude larger than the previously estimated value. However, the present CO data do not in themselves allow us to determine the origin of cold CO at the systemic velocity. Accordingly, the $\\rm{H_2}$ column density of $1 \\times 10^{22}$ $\\rm{cm^{-2}}$ derived above is a lower limit.  Similarly, the warm component at $\\sim$ 273 K, which is the dominant contributor to Figure~\\ref{fig:BoltzmanC6}, must be directly associated with the strongest absorption component, centered at $-$160 km~s$^{-1}$ in Figure~\\ref{fig:LineShape}. This component was detected in all of the transitions and is strongest at $J$=3--5, implying a gas temperature around 200 K. Outside of the Galactic center, the temperatures of molecular clouds do not exceed 100 K (Dickman \\& Clemens 1983), except in very limited volumes such as those close to the luminous stars and shocked regions. It thus seems unlikely this large column density of warm negative velocity gas is located outside of the nucleus of IRAS~08572$+$3915 NW. We propose that this component is close to the nucleus and is heated by the central engine of the AGN. Meijerink \\& Spaans (2005) showed that molecular clouds exposed to X-ray radiation (X-ray dissociation regions) can be heated to temperatures such as those observed here (see also Maloney et al. 1996; Meijerink, Spaans \\& Israel 2006). The warm molecular gas may also be collisionally heated due to interactions of portions of the gas having different velocities. However, although such turbulent heating is a possibility, the negative velocities of this warm component indicate that the bulk gas motion is outward from the nucleus.  The weak $+$100 km~s$^{-1}$ component seen in Figure~\\ref{fig:LineShape} does not contribute significantly to Figure~\\ref{fig:BoltzmanC6} because of its small column density. The equivalent widths of this component in the $P(6)$ and $P(16)$ line profiles are almost the same. From them we estimate that the temperature of the redshifted component is around 700 K and that the column density of CO is $\\sim10^{17}$ $\\rm{cm^{-2}}$. If the source of the heating of this cloud is the central engine, then the redshifted component would be expected to arise from clouds which lie closer to the central engine than those producing the blueshifted component. Outflows have been observed toward various types of AGN, but there are few examples of infall (e.g., Crenshaw et al. 2003; M\\\"uller S\\'anchez et al. 2009). We suggest that this velocity component traces the material fueling the central engine.   \\subsection{Densities and thicknesses of the warm clouds}\\label{sec:size}  The geometrical thickness of the absorbing cloud along the line of sight can be estimated from the column density derived above, if the gas particle density is known. As seen in the population diagram of Figure~\\ref{fig:BoltzmanC6}, the warm ($\\sim$ 270 K) component is well thermalized at least to the $J$=17 level. This indicates that the population distribution of the CO molecules in the rotational states is controlled by collisions with other species, presumably mostly $\\rm{H_2}$.  A high gas density is required to thermalize CO up to the $J$=17 level. The density at which collisional de-excitation competes with radiative de-excitation (the critical density; $n_c(\\rm{H_2})$) for CO $J\\to J-1$ is approximately $n_c(\\rm{H_2})=4\\times10^3$ $J^3$ $\\rm{cm^{-3}}$, as described in Kramer et al. (2004). The resulting value is $n_{c(17\\to16)}(\\rm{H_2})=2\\times10^7$ $\\rm{cm^{-3}}$. Collisional transitions with $\\Delta J >1$ are much less likely than $\\Delta J=1$ transitions and hence the critical density is approximately the above value. This lower limit to the actual density is extremely high compared to the densities of typical molecular clouds in the Galaxy ($n\\sim10^3$ $\\rm{cm^{-3}}$). Such densities are, however, found in post-shock regions of molecular clouds, where molecular outflows from young stars collide with ambient molecular gas (e.g., Burton et al. 1988; Vannier et al. 2001), and may also occur where molecular clouds collide.  Assuming that the absorbing cloud is homogenous, the thickness of the absorbing layer, $\\Delta d$, is \\begin{eqnarray} \\Delta d < \\frac{ N_{\\rm{H_2}} }{ n_c(\\rm{H_2}) }\\ \\approx \\ 4\\times10^{-4} \\ \\rm{pc}. \\end{eqnarray} The absorbing gas column thus appears to be extremely short along the line of sight. In general, dense molecular clouds are expected be highly clumpy, as a natural result of the interplay between the energy generated by self-gravity of the gas and radiative cooling. Clumpiness increases the size of molecular clouds with respect to the homogeneous case estimated above. It is rather difficult to estimate the volume filling factor, the fraction of absorbed volume in the beam area. Assuming that it is about 1--10\\%, similar to the values in highly clumped clouds in the Galaxy (Genzel et al. 1985), the thickness of the absorbing molecular clouds would be 10--100 times greater; i.e., 0.004--0.04 pc. This is still much less than the thicknesses of molecular tori assumed in previous studies of AGN (e.g., Jaffe et al. 2004; Wada \\& Tomisaka 2005).   \\subsection{Extinction estimates and locations of emitters and absorbers} \\label{sec:cont}  As described in Section~\\ref{sec:tmpcol}, we have detected the Pf$\\beta$ hydrogen recombination line. The intensity ratios of hydrogen recombination lines can be used to estimate the extinction by foreground dust, assuming ``Menzel case {\\it B}'' conditions (Osterbrock \\& Ferland 2005), for typical of H\\,{\\small II} regions. We thus compare the observed flux of Pf$\\beta$ (4.6538 $\\rm{\\mu m}$, $(3.2\\pm0.5)\\times10^{-17}$ $\\rm{W/m^2}$), with that of Br$\\gamma$ (2.1661 $\\rm{\\mu m}$, $(0.7\\pm0.2)\\times10^{-18}$ $\\rm{W/m^2}$) observed by Goldader et al. (1995). The flux ratio Pf$\\beta$/Br$\\gamma$ is $\\sim46$, whereas case {\\it B} recombination theory predicts a ratio of 0.58. The difference suggests a differential extinction of $E(M-K)=4.75$ mag to the recombination line emitting region. According to the standard extinction curve model for Milky Way dust (Draine 2003), this value corresponds to $A_V=57$ mag, implying a H column density of $N_{\\rm{H}}=1.07\\times10^{23}$ $\\rm{cm^{-2}}$ (using $N\\rm{(H)}/A_{V}=$ $1.9\\times10^{21}$ $\\rm{cm^2~mag^{-1}}$; Bohlin et al. 1978). The ratio could be considerably underestimated because the aperture used for our Pf$\\beta$ observation (0\\farcs60 $\\times$ 0\\farcs46) is 100 times smaller than that for Br$\\gamma$ (3\\farcs0 $\\times$ 9\\farcs0). Thus, the above column density is a lower limit. Case {\\it B} conditions are found in H\\,{\\small II} regions, but not in dense ionized winds; we caution that the physical conditions in the gas producing the hydrogen recombination lines is unknown.  From Section~\\ref{sec:partcoveragecase} the $\\rm{H_2}$ column density summed from all CO velocity components is $N_{\\rm{H_2}}\\sim4\\times10^{22}$ $\\rm{cm^{-2}}$, slightly lower than the lower limit estimated from the H recombination lines. However, this estimate of $N(\\rm{H_2}$) is itself a lower limit in that, as discussed earlier, some or all of the cold CO may reside in diffuse gas. X-ray observations toward IRAS~08572$+$3915 imply that the H column density to the AGN exceeds $N_{\\rm{H}}\\gg10^{24}$ $\\rm{cm^{-2}}$ (Risaliti et al. 2000). For comparison, the total hydrogen column density estimated from 9.7 $\\rm{\\mu m}$ silicate dust absorption is $N_{\\rm{H}}\\sim1\\times10^{23}$ $\\rm{cm^{-2}}$ ($\\tau_{\\rm{9.7}}/A_V=$ 0.05--0.1; Roche \\& Aitken 1984, 1985; $\\tau_{\\rm{9.7}}=4.2$; Spoon et al. 2006), comparable to the lower limit to the $\\rm{H_2}$ column density derived from the CO absorption. We conclude that the column density of hydrogen in front of the AGN, although highly uncertain, exceeds $10^{23}$ $\\rm{cm^{-2}}$.  Due to the large uncertainties in the above estimates, we cannot confidently assign relative locations along the line of site to the X-ray emitting region, recombination line emitting, infrared continuum emitting, and warm molecular absorbing regions based solely on estimates of foreground column densities of hydrogen. On the basis of energetics, however, one expects that the X-ray emitting gas must be interior to the recombination line emitting region. Both of those must be interior to the infrared continuum emitting regions at 5 and 10 microns, which themselves must be internal to the CO and silicate absorptions, respectively. These assignments are in fact consistent with the column densities derived above if they are actual values rather than (in some cases) lower limits. Also these suggest that the infrared continuum emission comes, not from the central engine itself, but from a location somewhat further from the engine. One expects that the continuum emission in the $M$-band originates close to the inner edge of the dusty molecular clouds closest to the AGN. The dust closest to the AGN probably has a high temperature, roughly equal to the sublimation temperature around $\\sim$ 1500 K. This hot dust produces thermal emission whose spectral energy distribution peaks at near-infrared wavelengths, and also contributes significantly at 5 $\\rm{\\mu m}$ continuum emission (Neugebauer et al. 1987; Barvainis 1992; Kobayashi et al. 1993). The CO absorption lines would then be observed largely against continuum emission produced by this hottest dust nearest to the AGN, while the silicate absorption would occur largely against dust located somewhat more distant from it.     Our spectrum toward IRAS~08572$+$3915 NW clearly separates the CO absorption into three velocity components; a cold component at the systemic velocity, a strong and warm component that is highly blueshifted ($-160$ km~s$^{-1}$), and a weaker, but much warmer component that is highly redshifted ($+$100 km~s$^{-1}$). From the population diagram of the observed CO absorption lines, we deduce that the blueshifted absorbing gas is dense ($> 10^{7}$ $\\rm{cm^{-3}}$) and warm ($\\sim$ 270 K) and has a CO column density of $N_{\\rm{CO}} \\sim 4.5 \\times 10^{18}$ $\\rm{cm^{-2}}$, which is equivalent to a $\\rm{H_2}$ column density of $N_{\\rm{H_2}} \\sim 2.5 \\times 10^{22}$ $\\rm{cm^{-2}}$. The redshifted absorbing gas appears to be produced by even warmer and presumably denser gas. These high temperatures and high densities, as well as the small column lengths ($<$ 0.04 pc), indicate that these absorption components are in close proximity to the central engine.  The physical properties that we derive are similar to the recent results from near- to mid-infrared spectroscopy of other ULIRGs. Lahuis et al. (2007) detected strong absorption bands of gaseous $\\rm{C_2H_2}$, HCN, and $\\rm{CO_2}$ toward deeply obscured (U)LIRG nuclei. Their data suggest the presence of warm (200--700 K) and dense ($> 10^7$ $\\rm{cm^{-3}}$) gas. Armus et al. (2007) argued that this gas occupies only a small fraction of the nuclear region ($\\sim$ 0.01 pc) near the mid-infrared continuum source. However, because all of their observations were made at much lower spectral resolution, they did not resolve individual velocity components. Our observations at high spectral resolution shed additional light on the geometrical structure and kinematics of the molecular clouds in one ULIRG and our conclusions may apply to others as well.  The extremely deep silicate absorption in some ULIRGs requires their nuclear sources to be deeply embedded in optically and geometrically thick material (Levenson et al. 2007). Such absorptions together with pointlike infrared nuclear sources often are taken as signposts of the existence of AGNs (Imanishi, Dudley \\& Maloney 2006). However, Spoon et al. (2006) discovered crystalline silicate features toward IRAS~08572$+$3915 in addition to the deep amorphous silicate feature, which can be a sign of high-temperature grain processing by the merger-triggered star-formation driven dust injection.  Various kinds of theoretical models for the material local to the AGN and obscuring it have been proposed. In all of them the central engines are obscured by dusty, optically and geometrically thick molecular clouds in toroidal structures; for example, a torus (Schartmann et al. 2005), a flared disk (Fritz et al. 2006), and discrete clouds (Elitzur \\& Shlosman 2006). Krolik \\& Begelman (1988) argued that the obscuring torus is composed of high-density molecular clouds. Nenkova et al. (2002) concluded that the obscuring torus is composed of a large number of small but dense molecular clouds. Indeed, the thermal dust emission from the torus has been spatially resolved by recent IR interferometric techniques and shown to have a clumpy or filamentary dust structure (Jaffe et al. 2004; Tristram et al. 2007). Additional evidence for the existence of the dust torus with many clumpy clouds comes from the detection of the silicate feature at $\\sim$10 $\\rm{\\mu m}$ in the spectral energy distribution of AGNs. The recent discoveries of silicate emission from QSOs (Siebenmorgen et al. 2005; Hao et al. 2005), Seyferts (Hao et al. 2007), and LINERs (Sturm et al. 2005) by Spitzer provide strong evidence for a clumpy torus model. From a theoretical point of view, according to the 3-dimensional radiative transfer model of clumpy dust tori of AGN (Vollmer, Beckert \\& Duschl 2004; Beckert \\& Duschl 2004; H$\\rm{\\ddot{o}}$nig et al. 2006; H$\\rm{\\ddot{o}}$nig \\& Beckert 2007; Nenkova et al. 2008), the properties of the individual self-gravitating clouds are speculated to be high density ($n \\sim 10^{7-8}$ $\\rm{cm^{-3}}$), large column density ($N_{\\rm{H}} \\sim 10^{24}$ $\\rm{cm^{-2}}$), and small volume filling factor ($\\sim$ 0.03). Wada (2007) investigated the possible origins of molecular absorption, based on a 3-dimensional hydrodynamic model of the inhomogeneous interstellar medium around the AGN. They demonstrated that the warm and dense molecular clouds, such as those discovered here via their CO absorption, can survive in the form of high-density clumps in the inhomogeneous interstellar matter around the AGN.  Our high resolution spectroscopy has successfully detected warm dense gas with complex velocity structure, containing both highly blueshifted and highly redshifted components. Moreover, the CO spectrum suggests that the absorbing clouds have a large area covering factor ($>0.6$) but a small volume filling factor (thickness $<$ 0.04 pc). Because the velocity dispersions of the redshifted and blueshifted components are so large, it seems inconceivable that each exists in a single narrow sheet; hence each component must itself be composed of numerous well-separated thin sheets of dense gas that are detached from the continuum source.  The redshifted and blueshifted components may trace the fueling of the central source and the ejection of material from it. However, most previous theoretical studies (e.g., Wada \\& Tomisaka 2005) have concluded that the dominant velocity structure in the torus is systematic rotation and turbulence. For IRAS~08572$+$3915 NW such models are apparently inconsistent with our observed multiple, discrete velocity components. The picture that the absorbing clouds in the torus rotate around the AGN cannot reproduce the observed CO absorptions in IRAS~08572$+$3915 NW for the following reasons: (1) Absorption by gas clouds in orbital motion around the AGN should reproduce a broad absorption profile centered at the systemic velocity, because it will be make up of the integrated absorption of the entire line-of-sight velocity. The observed deep absorption lines with clear line separation into the blue-shifted and red-shifted components cannot be expected; (2) The redshifted absorption is extremely weak relative to the blueshifted absorption; and (3) The redshifted gas has much higher temperature than the blueshifted gas.  We conclude that the velocity-shifted line components do not arise in a stable rotating torus-like gaseous structure, but are more likely associated primarily with ejection and injection of a dense and warm gas close to the AGN. Our results suggest that, at least in this case, the environment immediately surrounding the AGN does not appear to be as simple as that proposed in the unified scheme of AGNs.     We have observed and analyzed a high-resolution $M$-band spectrum of the obscured AGN, IRAS~08572$+$3915 NW, which contains strong absorption lines of the fundamental ($v$=1$\\leftarrow$0) band of CO. The main results and conclusions are as follows. \\\\ (1) Absorption lines of $^{12}\\rm{C}^{16}\\rm{O}$ $v$=1$\\leftarrow$0 are seen up to an excitation level of $J$=17. \\\\ (2) The absorbing gas has three velocity components: a cold component at the systemic velocity; a strong and warm component that is highly blueshifted ($-160$ km~s$^{-1}$); and a weaker, but much warmer component that is highly redshifted ($+$100 km~s$^{-1}$). \\\\ (3) Assuming LTE in the absorbing molecular gas, the temperatures of the systemic and blueshifted components are $23\\pm1$ K and $325\\pm5$ K, and the CO gas column densities are $N_{\\rm{CO,cold}}=(5.74\\pm0.06) \\times 10^{17}$ $\\rm{cm^{-2}}$ and $N_{\\rm{CO,warm}}=(2.65\\pm0.04) \\times 10^{18}$ $\\rm{cm^{-2}}$, respectively, for a covering factor of unity. For dense molecular gas these values correspond to $N_{\\rm{H_2,cold}}=(3.19\\pm0.03) \\times 10^{21}$ $\\rm{cm^{-2}}$ and $N_{\\rm{H_2,warm}}=(1.47\\pm0.02) \\times 10^{22}$ $\\rm{cm^{-2}}$. In the case of a covering factor of 0.6, which seems more likely based on optical depth arguments, the temperatures of these components are $24\\pm1$ K and $273\\pm2$ K, and the CO column densities are $N_{\\rm{CO,cold}}=(1.98\\pm0.01) \\times 10^{18}$ $\\rm{cm^{-2}}$ and $N_{\\rm{CO,warm}}=(4.48\\pm0.04) \\times 10^{18}$ $\\rm{cm^{-2}}$. These values correspond to $N_{\\rm{H_2,cold}}=(1.10\\pm0.01) \\times 10^{22}$ $\\rm{cm^{-2}}$ and $N_{\\rm{H_2,warm}}=(2.49\\pm0.02) \\times 10^{22}$ $\\rm{cm^{-2}}$, respectively. The redshifted component has a higher temperature ($\\sim$ 700 K) than the blueshifted component and its column density is considerably less ($N_{\\rm{CO,hot}} \\sim10^{17}$ $\\rm{cm^{-2}}$, yielding $N_{\\rm{H_2,hot}} \\sim10^{21}$ $\\rm{cm^{-2}}$). The H$_2$ column density associated with the cold gas at near systemic velocities may be considerably larger than the above values if the cold CO is largely in diffuse gas clouds. \\\\ (4) From an estimate of the critical density, we demonstrate that the molecular gas with warm temperature must be very dense ($n_{c}(\\rm{H_2}) > 2\\times10^7$ $\\rm{cm^{-3}}$) and geometrically thin ($<$ 0.04 pc) along the line of sight, even if it is highly clumpy. \\\\ (5) The presence of discrete, highly blueshifted and highly redshifted absorption components close to the central engine does not fit the standard unified model of AGNs.   \\bigskip  We acknowledge all of the staff and crew of the Subaru Telescope and NAOJ for their valuable assistance in obtaining this data and their continuous support for the construction and maintenance of IRCS. Particularly, we are grateful to H.~Terada and S.~Oya for technical support of these observations. We thank S.~Oyabu, S.~Matsuura, I.~Yamamura, K.~Wada, and T.~Kawaguchi for many fruitful discussions. T.~R.~G.'s research is supported by the Gemini Observatory, which is operated by the Association of Universities for Research in Astronomy, on behalf of the International Gemini Partnership of Argentina, Australia, Brazil, Canada, Chile, the United Kingdom, and the United States of America."
    },
    "1208/1208.5665_arXiv.txt": {
        "abstract": "\\myabstract \\newline \\begin{center} \\rule{38mm}{0.2mm} \\end{center}  \\myabstract ",
        "introduction": "Spatially and temporally coherent jets are a common feature of turbulent flows in planetary atmospheres  with the banded winds of the giant planets constituting a familiar  example  % \\citep{Vasavada-and-Showman-05}. \\cite{Fjortoft-1953} noted that the conservation of both energy and enstrophy in dissipationless barotropic flow implies that transfer of energy among spatial spectral components results in energy accumulating at the largest  scales. This argument provides a conceptual basis for understanding the observed tendency for formation  of large scale structure from small scale turbulence in  planetary atmospheres. However, the observed large scale structure is dominated by zonal jets with specific form and, moreover,  the scale of these jets is distinct from the largest scale  in the flow. \\cite{Rhines-1975}  argued that  the observed spatial scale of jets in beta-plane turbulence results from arrest of upscale energy transport at the  length scale, $\\sqrt{u/\\beta}$, where $\\beta$ is the meridional gradient of planetary vorticity and $u$ is the root mean square velocity in the turbulent fluid. In Rhines' interpretation this is the scale at which the turbulent energy cascade is intercepted by the formation of propagating Rossby waves. \\cite{Balk-etal-1991} extended Rhines' argument by showing  that in addition to energy and enstrophy, dissipationless barotropic turbulence conserves a  third quadratic  invariant, called zonostrophy, which constrains the large scale structures in dissipationless beta plane turbulence to be predominantly zonal (cf.~\\cite{Balk-etal-2009}). While these results establish a conceptual basis for expecting large scale zonal structures to form in beta plane turbulence, the physical mechanism of jet formation, the structure of the jets, and their dependence on parameters remain to be determined.  One mechanism for formation of jets is  vorticity  mixing resulting from Rossby wave breaking which leads to homogenization of vorticity in localized regions and formation of vorticity  staircases. The risers of these staircases correspond  to thin  prograde  jets located at the latitudes of steep vorticity  gradients  separating  parabolic retrograde jets corresponding to the well  mixed steps of the staircase \\citep{Baldwin-etal-2007, Dritchel-2008}. While vorticity staircases have been obtained in numerical simulations ~\\citep{Scott-Dritchel-2012}, in many cases mixing is insufficient to produce a staircase structure. Moreover, jets are observed to form from a bifurcation at infinitesimal perturbation amplitude and in the absence of wave breaking~\\citep{Farrell-Ioannou-2003-structural}.   Arguments based on equilibrium statistical mechanics  have  also been advanced to explain emergence  of jets e.g. by \\cite{Miller-1990} and \\cite{Robert-Sommeria-1991}. This theory is based on the principle that dissipationless turbulence tends to produce configurations that maximize entropy while conserving both energy and enstrophy. These maximum entropy configurations in beta plane barotropic turbulence assume the form, when observed at large scale, of zonal jets or  vortices (cf.~\\cite{Bouchet-Venaile-2012}). However, the relevance of these results to planetary flows that are strongly forced and dissipated and therefore out of equilibrium remains to be shown.       An important constraint on theories of jet maintenance is that the primary mechanism by which planetary turbulent jets are maintained is eddy momentum flux systematically directed up the mean velocity gradient; and this up-gradient momentum flux is produced by a broad spectrum of eddies, implying that the large-scale jets are maintained by spectrally nonlocal interaction between the eddy field and the large-scale zonal jets. This has been verified in observational studies on Jovian atmosphere~\\citep{Ingersoll-etal-2004,Salyk-etal-2006} and in numerical simulations~\\citep{Nozawa-and-Yoden-97,Huang-Robinson-98}. \\citet{Wordsworth-etal-2008} studied jet formation in rotating tanks and  found strong evidence confirming that jets are maintained by non-local energy transfer.  Laminar instability of a meridional Rossby wave or of a zonally varying meridional flow   can generate zonal fows \\citep{Lorenz-1972, Gill-1974,Manfroi-Young-99,  Berloff-etal-2009a, Connaughton-etal-2010}. Equations for the dynamics of these jets in the weakly nonlinear  limit were obtained by \\citet{Manfroi-Young-99}. This instability,  referred to as modulational instability, produces spectrally nonlocal transfer to the zonal flow  from the forced meridional waves but presumes a constant source of these finite amplitude meridional waves. In baroclinic flows, baroclinic instability has been advanced as the source of these coherent waves \\citep{Berloff-etal-2009a}.     Stochastic structural stability theory (SSST contracted to S3T) addresses turbulent jet dynamics as a two-way interaction between the mean flow and its consistent  field of turbulent eddies~\\citep{Farrell-Ioannou-2003-structural}. Both S3T and  modulational instability  involve non-local   interactions in wavenumber space  but these theories differ in that   in S3T the mean flow is supported by its interaction with a broad turbulence spectrum rather than with  specific  waves. In fact, S3T is a non-equilibrium statistical theory that provides a closure comprising a dynamics for the evolution of the mean flow together with its consistent field of eddies. In S3T the dynamics of the turbulence statistics required by this closure are obtained from a   stochastic turbulence model (STM), which provides accurate eddy statistics for the atmosphere at large scale~\\citep{Farrell-Ioannou-1993d, Farrell-Ioannou-1994a,Farrell-Ioannou-1995,% Zhang-Held-99}. % \\citet{Marston-etal-2008} have shown that the S3T system is obtained by truncating the infinite hierarchy of cumulant expansions to second order and they refer to  the S3T system as the second order cumulant expansion (CE2). In S3T, jets initially arise as a linear instability of the interaction between an infinitesimal jet perturbation and the associated eddy field and finite amplitude jets result from nonlinear equilibria continuing from these instabilities. Analysis of this jet formation instability determines the bifurcation structure of the jet formation process as a function of parameters. In addition to  jet formation bifurcations, S3T predicts  jet breakdown bifurcations as well as the structure of the emergent jets, the structure of the finite amplitude equilibrium jets they continue to, and the structure of the turbulence accompanying the jets. Moreover, S3T is a dynamics so it predicts   the time dependent trajectory of the statistical mean turbulent state as it evolves and, remarkably, the mean turbulent state is often predicted by S3T to be time dependent in the sense that the statistical mean state of the turbulence evolves  in a manner predicted by the theory \\citep{Farrell-Ioannou-2009-plasmas}. The formation of zonal jets in % planetary  turbulence was  studied as a bifurcation problem in S3T by \\cite{Farrell-Ioannou-2003-structural, Farrell-Ioannou-2007-structure, Farrell-Ioannou-2008-baroclinic, Farrell-Ioannou-2009-equatorial,  Farrell-Ioannou-2009-closure,Bakas-Ioannou-2011,  Srinivasan-Young-2012,Parker-Krommes-2013-generation}. A continuous formulation of S3T developed by  \\cite{Srinivasan-Young-2012} has facilitated analysis of the physical processes that give rise to the S3T instability and construction of  analytic expressions for the growth rates of  the S3T instability  in  homogeneous beta-plane turbulence \\citep{Srinivasan-Young-2012,Bakas-Ioannou-2013-jas}. Recently, the analogy between the dynamics of pattern formation   and zonal jet emergence  in the context of S3T was studied by~\\citet{Parker-Krommes-2013-generation}.              Relating S3T  to jet dynamics in fully nonlinear turbulence is facilitated by studying the quasi-linear model  which is intermediate between the nonlinear model and S3T. The quasi-linear (QL) approximation  to the full nonlinear dynamics (NL) results when eddy-eddy interactions are not explicitly included in the dynamics but are either neglected entirely or replaced with a simple stochastic parameterization, so that no  turbulent cascade occurs in the equations for the eddies, while interaction between the eddies and the zonal mean flow is  retained fully  in the zonal mean equation. S3T is essentially QL with the additional assumption of an infinite ensemble of eddies replacing the single realization evolved under QL. Although the dynamics of S3T and QL are essentially the same, by making the approximation of an infinite ensemble of eddies, the S3T equations  provide  an autonomous and fluctuation-free dynamics of the statistical mean turbulent state,  which transforms QL from a simulation of turbulence into a predictive theory of turbulence.    A fundamental attribute of QL/S3T is that the nonlinear eddy-eddy cascade of NL is suppressed in these systems. It follows that agreement in predictions of jet formation and equilibration  between NL and QL/S3T provides compelling evidence that cascades are not required for jet formation and theoretical support for observations showing that the turbulent transfers of momentum maintaining finite amplitude jets are non-local in spectral space.    Previous studies demonstrated that unstable jets maintained by mean flow body forcing can be equilibrated using QL dynamics \\citep{Schoeberl-Lindzen-84,DelSole-Farrell-1996,OGorman-Schneider-2007,Marston-etal-2008}. In contrast to these studies, in this work we investigate  the spontaneous emergence and equilibration of jets from homogeneous turbulence in the absence of any coherent external forcing at the jet scale.  S3T predicts that infinitesimal perturbations with zonal jet form organize homogeneous turbulence to  produce systematic up-gradient fluxes giving rise to exponential  jet growth and eventually to the establishment  of finite amplitude equilibrium jets.  Specifically,  the S3T  equations  predict initial formation of jets by the most unstable eigenmode of the linearized S3T dynamics. In agreement with S3T, \\cite{Srinivasan-Young-2012} found  that their NL simulations exhibit jet emergence  from a homogeneous turbulent state  with subsequent establishment of finite amplitude jets, while noting quantitative differences between bifurcation parameter values predicted by S3T and the parameter values for which jets were observed to emerge  in NL.  \\cite{Tobias-Marston-2013} also investigated the correspondence of CE2 simulations of jet formation with corresponding NL simulations and found that CE2 reproduces the jet structure, although they noted some differences in the second cumulant, and suggested a remedy by inclusion of higher cumulants.  In this paper we  use NL and its  QL counterpart together with S3T to examine further the dynamics  of emergence and equilibration of jets  from turbulence. Qualitative agreement in bifurcation behavior among these systems, which is obtained for all the spatial turbulence forcing distributions studied, confirms that the S3T instability mechanism is responsible for the formation and equilibration of jets.  Quantitative agreement is obtained for bifurcation parameters between NL and QL/S3T when account is taken of the modification of the turbulent spectrum that occurs in NL but not in QL/S3T.  Remarkably,  a primary component of this spectral modification can itself be traced to S3T instability, but of non-zonal rather than of zonal form.  We investigate the formation and equilibration of these non-zonal S3T instabilities and the effect these structures have on the equilibrium spectrum of beta-plane turbulence.  We also investigate circumstances under which non-zonal structures are modified and suppressed by the formation of zonal jets.  A dynamic of potential importance to climate is the possibility of multiple equilibria of the statistical mean turbulent state being supported with the same system parameters \\citep{Farrell-Ioannou-2003-structural,Farrell-Ioannou-2007-structure,Parker-Krommes-2013-generation}. We verify existence of  multiple equilibria, predicted by S3T, in our NL simulations. Finally, we show that weak jets result from stochastic excitation by the turbulence of stable S3T modes, which demonstrates the physical reality of the stable S3T modes. Turbulent fluctuation induced excitation of these weak local jets and the weak but zonally extended jets that form at slight supercriticality in the jet instability bifurcation may explain the enigmatic latent jets of \\cite{Berloff-etal-2011}. ",
        "conclusions": "In this work  predictions of  S3T for  jet formation and equilibration in barotropic beta-plane turbulence were critically compared with results obtained using QL and NL simulations. The qualitative bifurcation structure  predicted by S3T for emergence of zonal jets from a homogeneous turbulent state was confirmed  by both  the QL and NL simulations. Moreover, the finite amplitude  equilibrium jets in NL and QL  simulations were found to be as predicted by the fixed point solutions of  S3T. Differences in jet formation bifurcation parameter values between NL and QL/S3T were reconciled by taking account of the fact that the spectrum of turbulence is substantially modified in NL.  Remarkably, the modification of the spectrum in NL could be traced in large part to emergence of non-zonal  structures through S3T instability. When account is taken of the modification of the turbulent spectrum resulting substantially from these non-zonal  structures, S3T also provides quantitative agreement with the threshold values for the emergence of jets in NL. The influence of the background eddy spectrum on the S3T dynamics was found to be immediate, in the sense that  in spin-up simulations jets emerge in accordance with the instability calculated on the temporally developing spectrum. The fact that jets are prominent in observations is consistent with the robust result that when a jet structure emerges it has primacy over the non-zonal structures, so that even if the jet eigenfunction is not the most linearly S3T unstable eigenfunction, the jet still emerges at finite amplitude as the dominant structure.       These results confirm that jet emergence and equilibration in barotropic beta-plane turbulence results from the cooperative quasi-linear mean flow/eddy instability that is  predicted by S3T. These results also establish that turbulent cascades are not required for the formation of zonal jets in beta-plane turbulence. Moreover, the physical reality of the manifold of  stable modes arising  from cooperative interaction between incoherent turbulence and coherent jets, which is predicted by S3T,  was verified in this work by relating  observations of intermittent jets in NL and QL to stochastic excitation by the turbulence of  this manifold of  stable S3T modes.   S3T provides an autonomous, deterministic nonlinear closure of turbulence dynamics at second order that  provides an attractive vehicle for further investigation of the dynamics of  turbulent flows.           \\begin{acknowledgment} The authors would like to  acknowledge discussions with N. Bakas, F. Bouchet, K. Srinivasan and W. Young. Navid Constantinou acknowledges the support of  the Alexander S. Onassis Public Benefit Foundation. Brian Farrell was supported by  NSF AGS-1246929 and ATM-0736022. Brian Farrell and Petros Ioannou acknowledge the  hospitality during June 2012 of the  Aspen Center for Physics (supported by NSF under grant No.~1066293) where part of this paper was written. Petros Ioannou acknowledges  the generous support of the John S. Latsis Foundation under  ``Research Projects 2011\". \\end{acknowledgment}         \\ifthenelse{\\boolean{dc}} {} {\\clearpage}   \\begin{appendix}[A]% \\addcontentsline{toc}{section}{\\protect\\numberline{}Appendices}"
    },
    "1208/1208.5386_arXiv.txt": {
        "abstract": "We estimate the \\HI intensity fluctuation power spectrum for a sample of 18 spiral galaxies chosen from  THINGS. Our analysis spans a large range of length-scales from $\\sim 300 \\, {\\rm pc}$ to $\\sim 16 \\, {\\rm kpc}$ across the entire galaxy sample.  We find that the power spectrum of each galaxy  can be well fitted by a  power law $P_{\\rm HI}(U) = A\\ U^{\\alpha}$,  with  an index $\\alpha$ that varies from galaxy to galaxy. For some of the galaxies the  scale-invariant power-law  power  spectrum  extends to length-scales that are comparable to the size of the galaxy's disk. The distribution of $\\alpha$ is strongly peaked with $50 \\%$ of the values in the range $\\alpha=-1.9$ to $1.5$, and a mean  and  standard   deviation of  $-1.3$ and $0.5$ respectively. We find no significant correlation between $\\alpha$ and  the star formation rate, dynamical mass, \\HI mass or velocity dispersion of the galaxies.  Several earlier studies that have measured the power spectrum within our Galaxy on length-scales that are considerably smaller than $500 \\, {\\rm pc}$  have  found  a power-law power spectrum with $\\alpha$ in the range $\\approx -2.8$ to $-2.5$. We propose a picture where we interpret the values in the range $\\approx -2.8$ to $-2.5$ as arising from three dimensional (3D) turbulence in the Interstellar Medium (ISM) on length-scales smaller than the galaxy's scale-height, and we interpret the values in the range $\\approx -1.9$ to $-1.5$ measured in this paper  as arising from two-dimensional ISM turbulence in the plane of the galaxy's disk.  It however still remains a difficulty to explain the small galaxy to galaxy variations in the values of $\\alpha$ measured here. ",
        "introduction": "Scale-invariant  structures seen in galaxies are believed to be the result of  compressible turbulence. The source of the turbulence could be protostellar winds, supernova, rotational shear, spiral arm shocks, etc. (see \\citealt{2004ARA&A..42..211E} for a review). A variety of statistical measures have been proposed to characterize the structures seen in the ISM (see \\citealt{2000ApJ...537..720L, 2001ApJ...555L..33P, 2002ApJ...566..289B, 2004ApJ...616..943L, 2010ApJ...708.1204B}). Of these, power spectrum is most popular. In the case of compressible  turbulence   the power spectrum is expected to have a power law shape, and the slope $\\alpha$ of the power law contains information as to the nature of the turbulence.  In the case of the galactic neutral hydrogen (\\HI) power spectrum analysis was first done by \\citet{1983A&A...122..282C} who found that the slope $\\alpha$ of the power is $\\sim -2.7$ for spatial scales of $\\sim 5$ pc to $\\sim 100$ pc. Similar results were found by \\citet{1993MNRAS.262..327G} for slightly  larger length scales, i.e. $\\sim 200$ pc. The very small scales have been probed using absorption studies against extended sources. Power law slopes in the range $\\sim -2.5$ to  $-2.8$ have been found on scales of $0.01$ pc to $3.0$ pc by \\citet{2000ApJ...543..227D} and \\citet{2010MNRAS.404L..45R}. Studies of the \\HI in the LMC \\citep{1999MNRAS.302..417S} and SMC \\citep{2001ApJ...548..749E} also show that the intensity fluctuations have a power law spectrum.  Recently, \\citet{2006MNRAS.372L..33B} [henceforth Paper I] presented a visibility-based formalism for determining the power spectrum of external galaxies with extremely weak \\HI emission. Using this formalism they have found that the power spectrum of the nearby faint  (M$_{\\rm B} \\sim -10.9$) dwarf galaxy DDO~210 is a power law, i.e, $P(U) = A\\ U^{\\alpha}$ with  $\\alpha \\sim -2.6$. This is consistent with the values of $\\alpha$ measured for our Galaxy and  the nearby galaxies like LMC and  SMC.  In a subsequent paper \\citep{2008MNRAS.384L..34D}  [henceforth Paper II] we have  used the same visibility based formalism to measure the \\HI power spectrum  of the  spiral galaxy NGC~628.  The power spectrum was found to be a power law (on scales of $0.8$ to $8$ kpc), but the slope $\\alpha$ was found to be $\\sim \u22121.6$, i.e.  much flatter than that determined in the earlier studies. The earlier  studies  all probed much smaller length scales ($\\leq 500$ pc). \\citet{2008MNRAS.384L..34D} proposed that the difference arises because at large scales the measurements corresponds to two dimensional (2D) turbulence in the plane of the galaxy's disk whereas the earlier observations were all restricted to length scales smaller than the scale height (or thickness) of the galaxy's disk where three dimensional (3D) turbulence is expected. This picture was subsequently verified   \\citep{2009MNRAS.397L..60D} [henceforth  Paper III] in the galaxy NGC~1058  whose \\HI power spectrum was found to be well described by a broken power  law, with slope $\\alpha=-1.0 \\pm 0.2$ at large length scales ($1.5$ to $10.0$ kpc) and another with slope $\\alpha=-2.5\\pm 0.6$ at small length scales ($0.6$ to $1.5$ kpc).  The transition, which occurs  at the  length scale $1.5$ kpc, was used to  estimate  the  scale height of the galaxy's \\HI disk as $490 \\pm 90$ pc. In a more recent study, \\cite{2010ApJ...718L...1B}   have observed  a similar break in the power spectrum determined from  {\\em Spitzer} observations of the LMC and  measured the line of sight thickness of LMC to be in the  range $100-200$ pc. In \\citet{2009MNRAS.398..887D} [henceforth paper IV] we have extended our study of the \\HI power spectrum to a sample of five dwarf galaxies. Paper IV also presents simulations which substantiates our interpretation of the two different slopes in terms of 2D and 3D turbulence respectively. In \\citet{2010MNRAS.405L.102D}  [henceforth Paper V] we have studied the \\HI power spectrum of      a harassed galaxy from the Virgo cluster, where the power spectrum  slope of the outer part is different from that of the inner part  consistent with harassement being dominant at the outer parts.  In this paper we present the power spectrum of \\HI intensity fluctuations of  $18$ spiral galaxies drawn from the THINGS sample. THINGS, \u201cThe HI Nearby Galaxy Survey\u201d, is an \\HI 21-cm emission survey of 34 nearby galaxies carried out using the NRAO Very Large Array (VLA) \\citep{2008AJ....136.2563W}\\footnote{We are indebted to Fabian Walter for providing us with the  calibrated \\HI data from the THINGS survey.} with an aim to investigate the nature of the interstellar medium (ISM), galaxy morphology, star formation and mass distribution across the Hubble sequence. \\citet{2008AJ....136.2648D}  present high resolution rotation curves for the galaxies in the THINGS sample. Star formation rate and efficiency of the THINGS galaxies are also extensively studied in \\citet{2008AJ....136.2846B, 2008AJ....136.2782L}. This is  as an ideal data set for a comparative study of  the scale invariant fluctuations observed  in the neutral ISM of  spiral galaxies. The rest of the paper is organized as follows. In Section~\\ref{sec:data} we discuss  the galaxies in our sample and briefly mention the power spectrum estimator used in this study. The results are  presented and discussed in Section~\\ref{sec:results}. Finally we conclude with our main results in the Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} We have measured  the angular power spectrum of 21-cm specific intensity fluctuations for a sample  of 18 spiral galaxies drown from THINGS.  This is the first comprehensive analysis of the  power spectra of  a moderately sized sample of external spiral galaxies.  For all the galaxies the estimated  power spectrum   can be well fitted with a single power law across a reasonably large  dynamical  range.  The power-law power spectrum, we find,  extends to  length-scales as large as  $\\sim$ 10 kpc indicating  the presence of scale-free structures in the ISM on length-scales that are comparable to the size of the galaxy.  In this analysis we have been able to measure the \\HI power spectrum over a large range of length-scales spanning from $\\sim \\ 300$ pc to $\\sim \\ 16$ kpc across the entire galaxy sample.  The power spectra, we find,  are well fit by power laws indicating the presence of scale-invariant fluctuations. Fifty percent of the galaxies in our sample have a measured power law index (slope) $\\alpha$  in the range  $-1.9$ to  $-1.5$ (Figure~\\ref{fig:hist}).   Only one galaxy, NGC~5457, has  a slope $\\alpha=-2.2$ which is steeper than  $-2$. All the other galaxies in our sample have slopes $\\alpha \\ge -2$.   A large number of earlier studies (summarized in \\citealt{2004ARA&A..42..211E}),   which have probed  small length-scales ranging from $~10$ pc to $~200$ pc, find a power law power spectrum with slope  $\\alpha \\approx-2.8$. This is believed to be the outcome of three dimensional (3D) compressible turbulence in the ISM. If the same  process extends to  length-scales larger than the galaxy's scale height, we would expect to see two-dimensional turbulence in the plane of the galaxy's disc.   Dimensional arguments given in Papers II and III   lead us to expect the slope of the 2D density fluctuations to be $\\alpha \\approx-1.8$. This, within estimated measurement errors,  is consistent with the slope that we have measured here for most of the galaxies in our sample. This  prompts us to believe that the small scales and large scale fluctuations  may both originate from the same physical process, presumably turbulence.  Energy input from supernova is believed to be the major driving mechanism for turbulence at small length-scales. However, it is unlikely that this mechanism would be able to generate large scale coherent structures as seen here. Further, the range of length scales that we probe here ranges from $\\sim 300$ pc to $~16$ kpc which has practically  no  overlap with the ranges of length scale  at  which turbulence is believed  to be operational in our Galaxy.  Our analysis   confirms that the power-law index  index has no correlation with  inclination angle, SFR, dynamical or total \\HI   mass of the galaxy.  At present we do not have any understanding of the  physical process responsible for the scale-invariant large-scale fluctuations measured here. The  mass fluctuations  in the galaxy's dark matter halo  is an interesting possibility that we plan to pursue in future.  Numerical studies of the magneto-hydrodynamic (MHD) turbulence \\citep{2005ApJ...624L..93B,2007ApJ...658..423K,2011ApJ...736...60T,2012arXiv1205.3792B} demonstrate that the Mach number of the medium is correlated with the  density fluctuation power spectrum index. Alternatively, the Mach number can also be estimated from the turbulent velocity dispersion of the medium for such MHD turbulence. We are presently investigating possibilities to separate the turbulent velocity dispersion in the \\HI gas from it's thermal counterpart. It may be possible that the \\HI column density as well as the velocity dispersion provides measure of similar quantities on the turbulent \\HI. However, it is not straightforward to carry over the results of the simulations performed for MHD turbulence to the turbulence in the \\HI gas, which may be characteristically different. This can be tested in future studies and give rise to a better understanding of the ISM turbulence in general."
    },
    "1208/1208.2273_arXiv.txt": {
        "abstract": "We present 248 precise Doppler measurements of Barnard's Star (Gl 699), the second nearest star system to Earth, obtained from Lick and Keck Observatories during 25 years between 1987 and 2012. The early precision was 20 \\ms{} but was 2 \\ms{} during the last 8 years, constituting the most extensive and sensitive search for Doppler signatures of planets around this stellar neighbor. We carefully analyze the 136 Keck radial velocities spanning 8 years by first applying a periodogram analysis to search for nearly circular orbits. We find no significant periodic Doppler signals with amplitudes above $\\sim$2 \\ms{}, setting firm upper limits on the minimum mass (\\msini) of any planets with orbital periods from 0.1 to 1000 days. Using a Monte Carlo analysis for circular orbits, we determine that planetary companions to Barnard's Star with masses above 2 \\mearth{} and periods below 10 days would have been detected. Planets with periods up to 2 years and masses above 10 \\mearth{} (0.0 3 \\mjup) are also ruled out. A similar analysis allowing for eccentric orbits yields comparable mass limits. The habitable zone of Barnard's Star appears to be devoid of roughly Earth-mass planets or larger, save for face-on orbits. Previous claims of planets around the star by van de Kamp are strongly refuted. The radial velocity of Barnard's Star increases with time at $4.515\\pm0.002$ \\msy{}, consistent with the predicted geometrical effect, secular acceleration, that exchanges transverse for radial components of velocity. ",
        "introduction": "\\label{intro}  To date, over 700 exoplanets have been identified orbiting other stars \\citep{Marcy2008,Mayor2011, Wright2011}, and another 2300 exoplanet candidates have been found from the Kepler spaceborne telescope \\citep{Batalha2012}, the majority of which are real planets \\citep{Morton2011, Lissauer2012}. Hundreds of exoplanets have now been discovered within 50 pc, most by precision Doppler surveys \\citep{Wright2011}. These nearest exoplanets provide the best opportunities for follow-up observations by the next generation of planet detection techniques, which now include numerous strategies, both ground- and space-based, such as direct imaging \\citep{Marois2008}, transit (e.g., \\citealt{Irwin2009, Muirhead2012b, Berta2012}), IR thermal signatures (e.g., \\citealt{Charbonneau2005}), and astrometry \\citep{Anglada2012}.  Searching the nearest stars for planets presents special challenges. These campaigns require large telescopes to conduct exhaustive long-term radial velocity (RV) surveys, and the very closest stars---those within a few pc---are mostly faint M dwarfs. While nearly 300 M dwarfs are currently being monitored for exoplanets \\citep{Johnson2010b, Delfosse2012}, relatively little radial velocity data on them were available until recently. The first planet orbiting an M dwarf was discovered in 2001 around Gl 876 \\citep{Marcy2001}. This M4V star has since been found to host four companions, including a 7.5 Earth-mass planet \\citep{Rivera2005, Rivera2010}. In the last few years, many planets have been found around other M dwarfs, including: Gl 832 (M1.5, \\citealt{Bailey2009}), Gl 649 (M2, \\citealt{Johnson2010b}), Gl 179 (M3.5, \\citealt{Howard2010}), HIP 12961 (M0, \\citealt{Forveille2011}), Gl 676 A (M0, \\citealt{Forveille2011}) Gl 433 (M1.5 \\citealt{Bonfils2011}), and Gl 667 C (M1, \\citealt{Bonfils2011, Delfosse2012}), increasing the number of currently known planetary companions around M dwarfs to 25 \\citep{Wright2011} at the time of writing.  \\cite{Johnson2007a} and \\cite{Johnson2010c} found a positive correlation between the frequency of jovian planets and host star mass, lending support to the core accretion model of planet formation (e.g., \\citealt{Kennedy2008}). It has been well-established that Jovian planets appear to form less frequently around M dwarfs than more massive stars \\citep{Johnson2010b}. Currently, the best estimates for occurrence rate of planets with $M_{\\rm P} \\sin i > 0.3$ \\mjup{} in orbits within 2.5 AU of their parent stars is $3.4^{+2.2}_{-0.9}\\%$ for stars with $M_{\\rm S} < 0.6$ \\msun{}, where $M_{\\rm P}$ and $M_{\\rm S}$ refer to the masses of the planet and the star, respectively, compared to $\\sim8\\%$ for F,G, and K stars \\citep{Cumming2008, Johnson2010b, Bonfils2011}. Recent works from both transit and RV surveys revealed that low-mass planets, rather than gas giants, are common around M dwarfs \\citep{Bonfils2011, Howard2011}. Surveys with a long time baseline and high precision such as this work are necessary for the detection of these low-mass planets.  Nearby stars with high proper motions exhibit changes in their RVs over time due to {\\it secular acceleration} \\citep{Stumpff1985}, an effect just at the limit of detectability for most surveys \\citep{Kurster2003}. We remove the effect of secular acceleration from our RVs and search these data for signals due to exoplanets. We use a Monte Carlo approach to place upper limits on the minimum mass of possible exoplanets. Finally we compare these observations to previous claims of planetary companions around Barnard's Star. ",
        "conclusions": "We have established firm upper limits to the minimum masses (\\msini{}) of planets around Barnard's Star for orbital periods ranging from a few hours to 20 years. For orbital periods under 10 days, planets with \\msini{} greater than two Earth masses would have been detected, but were not seen. For orbital periods under 100 days, planets with minimum masses under $\\sim3$ \\mearth{} would have been detected, but none was found. For periods under 2 years, planets with minimum masses over 10 \\mearth{} are similarly ruled out.  The two planets claimed by Peter van de Kamp are extremely unlikely by these 25 years of precise RVs. We frankly pursued this quarter-century program of precise RVs for Barnard's Star with the goal of examining anew the existence of these historic planets. Indeed, Peter van de Kamp remains one of the most respected astrometrists of all time for his observational care, persistence, and ingenuity. But there can be little doubt now that van de Kamp's two putative planets do not exist.  Even van de Kamp's model of a single-planet having 1.6 \\mjup{} orbiting at 4.4 AU \\citep{vandeKamp1963} can be securely ruled out. The RVs from the Lick and Keck Observatories that impose limits on the stellar reflex velocity of only a few meters per second simply leave no possibility of Jupiter-mass planets within 5 AU, save for unlikely face-on orbits.  The lack of planets above a few Earth masses near Barnard's Star runs counter to the discoveries of numerous mini-Neptunes, with sizes and masses slightly above those of Earth, found recently around M dwarfs. A detailed analysis of the planet candidates from the NASA Kepler mission shows an increasing number of small planets (2--4 R$_{\\oplus}$) around stars of decreasing mass, including the M dwarfs \\citep{Howard2011}.  \\cite{Howard2011} determined occurrence rates for planets with orbital periods less than 50 days. For planets of 2--4 R$_{\\oplus}$, the occurrence is 10\\% for G-type stars. But the occurrence of such low-mass planets linearly increases with decreasing $T_{\\rm eff}$, reaching seven times more abundant around cool stars (3600--4100 K) than around the hottest stars in the Kepler sample (6600--7100 K). Thus Kepler finds a large occurrence of 2--4 R$_{\\oplus}$ planets close-in to M dwarfs, just where our RV survey of Barnard's Star is most sensitive to Earth-mass planets. Yet, we found no planetary companions around Barnard's Star.  Similarly, the HARPS survey for M dwarfs has revealed numerous planets with \\msini{} of a few Earth masses around M dwarfs \\citep{Bonfils2011}. They examined 102 M dwarfs and found nine ``super-Earths,\" with two within the habitable zones of Gliese 581 and Gliese 667C. Extrapolating, they found that the occurrence of ``super-Earths\" in the habitable zone is $\\sim41$\\% for M dwarfs.  Thus, we have a lovely moment in science.  Two completely different planet-hunting techniques, Doppler measurements by HARPS to detect the reflex motion of stars, and brightness measurements by Kepler to detect the transits of planets, give similar and extraordinary results. Small planets, slightly larger or more massive than Earth, are apparently common around M dwarfs.  In contrast, the non-detection of planets above a few Earth masses around Barnard's Star remains remarkable as the detection limits here are as tight or tighter than was possible for the Kepler and HARPS surveys. The lack of planetary companions around Barnard's Star is interesting because of its low metallicity. This non-detection of nearly Earth-mass planets around Barnard's Star is surely unfortunate, as its distance of only 1.8 parsecs would render any Earth-size planets valuable targets for imaging and spectroscopy, as well as compelling destinations for robotic probes by the end of the century."
    },
    "1208/1208.2790_arXiv.txt": {
        "abstract": "In the theories of generalized modified gravity, the acceleration equation is generally fourth order. So it is hard to analyze the evolution of the Universe. In this paper, we present a class of generalized modified gravity theories which have the acceleration equation of second order derivative. Then both the cosmic evolution and the weak-field limit of the theories are easily investigated. We find that not only the Big-bang singularity problem but also the current cosmic acceleration problem could be easily dealt with. ",
        "introduction": "In history, the motivation for modifying GR (General Relativity) mainly comes from the fact that GR is not renormalizable. So it can not be conventionally quantized. In the first place, Utiyama and DeWitt showed that the renormalization at one-loop requires the higher order curvature terms in the action of gravity theories \\cite{uti:62}. Secondly, Stelle showed the corresponding gravity theories with these higher order terms are indeed renormalizable \\cite{ste:77}. Finally, when quantum effects or string theory are taken into account, the effective low energy gravitational action also requires higher order curvature invariants \\cite{bir:82,buc:92,vil:92}. So it was generally believed that the modifications to GR would be important only at the scales of very close to the Planck lengthy or Planck energy. Consequently, both the Big-bang singularity and black hole singularity are expected to be absent in the modified gravity theories \\cite{sta:80,bra:92,bra:93,muk:92,sha:90,tro:93}. This is the belief before 1998.  However, with the discovery of cosmic acceleration in 1998 \\cite{per:99,rie:98}, one realize that GR may also need to be modified on very large scale or at very low energy (or very weak gravitational field). These constitute the infrared modifications to GR, for example, the GDP (Dvali-Gabadadze-Porrati) model \\cite{dva:00}, the $1/R$ modified gravity model \\cite{car:04} and so on. Here we shall not produce an exhaustive list of references on modified gravity, but we prefer the readers to read the nice review paper by Sotiriou and Faraoni \\cite{faraoni:08} and the references~therein. In general, the equations of motion for the generalized modified gravity are of fourth order and one can expect the particle content of the theory would have eight degrees of freedom: two for the massless graviton, one in a scalar excitation and five in a ghost-like massive spin two field \\cite{hin:95}. The presence of ghost leads one to accept unphysical negative energy states in the theory and the property of unitarity is lost \\cite{haw:02}. This ghost problem is closely related to the higher order property of the theories.  So the purpose of this paper is to seek for the \\emph{second order} theories of gravity, at least in the background of spatially flat FRW (Friedmann-Robertson-Walker) Universe. Except for satisfying the requirement of second-order, the theories also meet ghost free conditions. Due to the property of second order of acceleration equation, the resulting Friedmann equation remains first order and the cosmic evolution of the universe is easily deal with.  The paper is organized as follows. In section II, we briefly review the generalized modified gravity theories. The equations of motion are presented. In section III, we propose the Lagrangian for the generalized modified gravity which are both ghost free and second order (in the background of spatially flat FRW Universe). In section IV, we investigate the cosmic evolution of some specific models of Lagrangian. In section V, we investigate the weak field limit of these models. Section IV gives the conclusion and discussion.  We shall use the system of units in which $G=c=\\hbar=1$ and the metric signature $(-,\\ +,\\ +,\\ +)$ throughout the paper. ",
        "conclusions": "In the theories of generalized modified gravity, the acceleration equation is generally fourth order. So it is hard to analyze the evolution of the Universe \\cite{car:04}. On the other hand, these theories are also plagued with the ghost problem. So the property of unitary of the theory is lost \\cite{haw:02}. In view of this point, we present a class of generalized modified gravity theories which have the acceleration equation of second order derivative and ghost free. Then we explore some specific examples for the Lagrangian function. We find both the cosmic evolution and the weak-field limit of the theories are easily investigated. Furthermore, not only the Big-bang singularity problem but also the current cosmic acceleration problem could be easily dealt with."
    },
    "1208/1208.2036.txt": {
        "abstract": "We report the results of a spectroscopic study of the high-mass protostellar object NGC 7538 IRS 9 and compare our observations to published data on the nearby object NGC 7538 IRS 1.  Both objects originated in the same molecular cloud and appear to be at different points in their evolutionary histories, offering an unusual opportunity to study the temporal evolution of envelope chemistry in objects sharing a presumably identical starting composition.  Observations were made with the Texas Echelon Cross Echelle Spectrograph (TEXES), a sensitive, high spectral resolution ($R = {\\lambda}/{\\Delta}{\\lambda}\\simeq$ 100,000) mid-infrared grating spectrometer.  Forty-six individual lines in vibrational modes of the molecules C$_2$H$_2$, CH$_4$, HCN, NH$_3$ and CO were detected, including two isotopologues ($^{13}$CO, $^{12}$C$^{18}$O) and one combination mode (${\\nu}_{4}+{\\nu}_{5}$ C$_2$H$_2$).  Fitting synthetic spectra to the data yielded the Doppler shift, excitation temperature, Doppler $b$ parameter, column density and covering factor for each molecule observed; we also computed column density upper limits for lines and species not detected, such as HNCO and OCS.  We find differences among spectra of the two objects likely attributable to their differing radiation and thermal environments.  Temperatures and column densities for the two objects are generally consistent, while the larger line widths toward IRS 9 result in less saturated lines than those toward IRS 1.  Finally, we compute an upper limit on the size of the continuum-emitting region ($\\sim$2000 AU) and use this constraint and our spectroscopy results to construct a schematic model of IRS 9. ",
        "introduction": "}  Details of the pre-Main Sequence (pre-MS) life cycles of low-mass stars such as the Sun ($M$ $\\sim$ 1$M_{\\odot}$) have been gleaned from decades of research; see \\citet{McKeeOstriker07} for a recent review.  This is thanks in part to the relative amenability of low-mass protostellar objects to study.  However, details of high-mass ($M$ $>$ 8$M_{\\odot}$) star formation are not as well understood at present, in part because high-mass stars are not generally forming near the solar neighborhood and remain deeply embedded during their pre-MS evolution.   An important unresolved issue is whether high-mass stars form in a manner resembling a scaled-up version of low-mass  star formation \\citep{KetoZhang10, Johnston11} or through other mechanisms uncharacteristic of the low-mass case such as collisions/mergers \\citep{Bonnell98, Bonnell05, BallyZinnecker05} and competitive accretion in clusters  \\citep{Bonnell01, Bonnell04}.  Clearly, the environments in which high-mass star formation takes place seem to influence the formation process itself.  Massive stars are most often seen to form in large groups (``OB associations'') near the edges of dense clumps in giant molecular clouds (GMCs; \\citealt{YZ07}).  \\citet {ElmegreenLada77} first suggested that spatially distinct subgroups of stars within OB associations trigger waves of self-propagating star formation in GMCs via shock/ionization fronts that progress from the outside in.  This process may account for observations of OB subgroups that lie in linear arrangements along the Galactic plane, showing a monotonically-increasing sequence of ages \\citep{Blaauw64,Blaauw91}. Rapid improvements in both observational capabilities and theoretical sophistication in recent years are advancing our knowledge of the circumstances of massive star formation and enabling these fundamental questions to be addressed.  The origin of these stars is of particular interest due to their role in the Galactic ecosystem.  In addition to serving as the principal sites of heavy element nucleosynthesis, the violent end of massive stars in supernovae injects considerable mechanical energy into the interstellar medium (ISM) while enriching it with metals \\citep{Arnett96}.  The debris are reconstituted in molecular clouds from which new stars form, often accompanied by circumstellar disks \\citep{Jiang08}.  These disks are evidently the source of accretion onto high-mass stars as they approach the MS in the Hertzsprung-Russell diagram and ignite hydrogen burning, then continue to accrete their final masses, moving parallel to the MS \\citep{YZ07}.  Relatively little is known about the structure and chemical composition of massive protostellar envelopes, and even less about how they evolve in time.  Ready comparison with envelope models for low-mass objects is instructive but only of limited use, in part due to the radically different radiation environment in the high-mass case.  The composition and structure of the envelope are of interest as they determine the final mass of the star through an accretion mechanism whose details are still unclear.  Piecing together the components of protostellar envelopes is complicated by the observation that most objects are deeply embedded, particularly early in the formation process.   Spectroscopy of molecules is useful in delineating physical structures at high extinction where other methods are limited.  Pure rotational transitions of molecules in the millimeter probe a sufficiently low energy regime but such observations are typically of low spatial resolution.  Better resolution is achieved in the infrared where many ro-vibrational molecular lines are located.  We present high resolution, mid-infrared spectra of a number of molecules toward an embedded infrared source in the NGC 7538 starforming complex, NGC 7538 IRS 9, and compare the results to a previous study of the nearby source NGC 7538 IRS 1 \\citep{Claudia}. Both sources are presumed to harbor high-mass protostellar objects in their interiors.  Their proximity to one another suggests that, having formed within the same molecular cloud, the starting chemical abundances for each object should have been substantially similar.  A comparison of their current states provides insights about differences in their ages and/or divergence in their individual evolutionary histories.  Weak free-free emission ($\\leq$ 60 $\\mu$Jy at 3.6 cm; \\citealt{Sandell05}) is observed toward IRS 9.  Further evidence for the comparatively primitive nature of the envelope of IRS 9 is given by the detection of various ices in its \\it Infrared Space Observatory \\rm (ISO) spectrum (\\citealt{Whittet96}; \\citealt{Gibb04}).  These ices are not seen toward IRS 1, which appears to be in a relatively advanced point in its pre-MS evolution.  It is a considerably stronger source at centimeter wavelengths and shows evidence of both ionized and molecular components in a bipolar outflow originating in a disk seen nearly edge-on \\citep{Sandell09, Qiu11}.  IRS 1 also shows signs of a circumstellar disk and an emerging compact \\ion{H}{2} region; it is thought to be a forming late O star undergoing active accretion, suggesting a more advanced evolutionary state \\citep {Lacy08, Beuther12}.  \\citet{Sandell09} further note that the system is heavily accreting ($\\dot{M}$ $\\sim$ 2 $\\times$ 10$^{-4}$ $M_{\\odot}$ yr$^{-1}$) and that the accretion may be episodic in nature.  On the basis of GHz line observations, \\citet{Surcis11} conclude that the accretion  appears to be the result of radial infall from a ``torus'' thousands of Astronomical Units (AU) in size, traced by H$_2$O and CH$_3$OH masers, instead of the Keplerian disk favored by \\citet{Pestalozzi04}.  Until recently, estimates of the distance to NGC 7538 placed it at $\\sim$3.0 kpc \\citep{HCS01, Brunt03, FR03, Balog04, RW05, Kameya06, Araya07}.  However, we adopted the distance of \\citet{Moscadelli09}, who find a value of 2.65$^{+0.12}_{-0.11}$ kpc from trigonometric parallax measurements.  The total luminosity of IRS 9 is about 3.5 $\\times$ 10$^{4}$ L$_{\\odot}$ (\\citealt{Sandell05}, corrected to the Moscadelli et al. distance) and it is a bright IRAS 12 $\\mu$m source at $\\sim$60 Jy \\citep{IRAS}.  A zero-age Main Sequence (ZAMS) star with this luminosity would have a spectral type of B0.5, an effective temperature $T_{eff}$ = 26200 K \\citep{Panagia73}, and a mass of $\\sim $13 M$_\\odot$ \\citep{Ekstrom12}.  \\citet{BoogertBlake04} suggest the existence of an inner, warm molecular emission region with a radius of $\\sim$70 AU; the dust temperature a tthis distance for an object of IRS 9's luminosity would be $\\sim$1000 K depending on the dust mass opacity.   Millimeter data show that IRS 9 is a site of active, massive star formation \\citep{Sandell05} in which young stellar objects (YSOs) are driving a set of bipolar molecular outflows on a dynamical timescale of $\\leq$ 20,000 years.  At least one of these outflows appears to be highly energetic.  \\citet{Mitchell91} report the discovery of an outflow in IRS 9 with a velocity of 110 km s$^{-1}$, which they interpret to have emerged as recently as 1,200 years ago, perhaps indicating the beginning of a recent episode of accretion.  Observations of these protostellar objects are  broadly consistent with the \\citet{ElmegreenLada77} picture of sequential high-mass star formation, as elaborated on by others in the particular case of NGC 7538 \\citep{Werner79, Fischer80, Dickel81}.  The \\citet {CampbellThompson84} model of NGC 7538 holds that a first generation of massive stars caused the visible \\ion{H}{2} region, which expanded and impinged on the neighboring molecular cloud; this compressed the cloud and generated shocks that caused the observed near-infrared fluorescent H$_2$ emission northwest of IRS 1. In this picture, high-mass star formation is progressing toward the southeast, and IRS 9 belongs to a second wave of star formation as it is placed further along the direction of propagation of the shock front.  These objects then represent two snapshots of the evolution of a massive young protostellar object in a particular cloud and are useful in addressing temporal evolution questions.  This paper is organized as follows.  In Section 2  we give details concerning how we obtained and reduced the data, followed by a detailed description in Section 3 of the molecular species detected.  The data analysis is described in Section 4 with emphasis on spectral modeling, and we give some interpretation to the results in Section 5.  We summarize these results and draw some conclusions in Section 6. ",
        "conclusions": "}  We have presented the method and results of a study of the embedded high-mass protostellar object NGC 7538 IRS 9 and compared our findings to a similar investigation of the related object NGC 7538 IRS 1.  We obtained high resolution, mid-infrared spectra of these objects in 46 ro-vibrational transitions of the fundamental bands of the molecules C$_2$H$_2$, CH$_4$, HCN, NH$_3$ and CO and a number of their isotopologues.  We also detected two lines of the ${\\nu}_4+{\\nu}_5$ combination mode of C$_2$H$_2$. From these observations we draw some broad conclusions.  1. IRS 9 appears to be a spatially-unresolved object on a scale of $\\sim$2000 AU.  2. With the exception of CO, whose lines are saturated, we did not observe $^{13}$C isotopologues of any organic molecules.    This may indicate a particularly large value of the $^{12}$C/$^{13}$C ratio, or the $^{13}$C isotopologue lines were simply below our threshold of detection.  3. There is no discernible, consistent trend in the abundance variations of C$_2$H$_2$, HCN, NH$_3$, and CH$_4$ with respect to either CO or H$_2$ between IRS 9 and IRS 1.  Column density variations between the two objects are an order of magnitude or less in each case.  4. The observation of gas-phase CS absorption toward IRS 1 and the evident non-detection of sulfur-bearing species like OCS in the gas phase toward IRS 9 may reflect a real S abundance variation between the two objects but can also be explained by different thermal conditions in each case.  A similar abundance discrepancy involving HNCO may indicate sufficiently low temperatures in the IRS 9 envelope to retain it as an ice and explain its detection in the solid phase by \\citet{Gibb04}.  If the abundances in all phases are actually comparable between the two sources, a relatively weak heating source is implied in the case of IRS 9.  While this is at odds with the somewhat higher inferred abundances of the other molecules toward the warmer environment of IRS 1, the suggested sightline in that direction through a disk atmosphere subjects those molecules to direct irradiation by protostellar UV.  This may drive photochemistry that lowers abundances of certain molecules as they are converted to other kinds.  5. However, the relative abundances of various molecular species with respect to CO and H$_2$ can also be explained by variations in the CO/H$_2$ ratio or the gas-to-dust ratio between the two objects.  Choosing between these interpretations would require better determining the CO and H$_2$ abundances toward both objects.  6. Lines of many species observed toward IRS 9 have higher Doppler $b$ values than those toward IRS 1, resulting in lines toward IRS 9 being less saturated than those toward IRS 1.  7. Given the non-detection of \\ion{Ne}{2} in our data and the continuum flux density at the expected wavenumber of the line (70$\\pm$10 Jy), we find a corresponding upper limit of $\\sim$10$^{45}$ photons s$^{-1}$ for the Lyman continuum flux in IRS 9.  This value is well below both the expected flux of a ZAMS star with the same luminosity and the ionizing flux of IRS 1 ($>$ 10$^{48}$ photons s$^{-1}$).  8. The observed pattern of emission and absorption components of ammonia lines in our spectra can be explained by a non-LTE radiative transfer effect involving ``pumping'' of lines on the $R$-branch by light from the 9.7 $\\mu$m silicate dust feature.  Predictions of the strength of this effect compare reasonably well with observed fluxes in $P$-branch lines.  9. Our data are consistent with a simple model of the IRS 9 system in which the sightline probes both a high-speed outflow and a quiescent envelope illuminated from behind by a ``dust photosphere'' such that line formation depth is wavelength-dependent.  The possibility of seeing the turbulent flow along the wall of the outflow cavity is supported by the broad, blueshifted wing of the C$_2$H$_2$ ${\\nu}_{4}+{\\nu}_{5}$ combination band that may arise in a shock.  A disk is included in the model given the system's inferred young age and the presence of a kinematically-indicated disk toward the (presumably) older IRS 1.  10. Observed differences in the abundance of various molecular species toward IRS 9 and IRS 1 imply an upper limit to the age of IRS 9 of $\\sim$ 10$^4$ yr, whereas IRS 1 may be an order of magnitude older.  This conclusion is consistent with the triggered star formation models of \\citet{ElmegreenLada77} as applied to NGC 7538 by \\citet{CampbellThompson84} and others."
    },
    "1208/1208.0871_arXiv.txt": {
        "abstract": "We demonstrate single-photon counting at 1550 nm with titanium-nitride (TiN) microwave kinetic inductance detectors. Energy resolution of  0.4 eV and arrival-time resolution of 1.2 microseconds are achieved. 0-, 1-, 2-photon events are resolved and shown to follow Poisson statistics. We find that the temperature-dependent frequency shift deviates from the Mattis-Bardeen theory, and the dissipation response shows a shorter decay time than the frequency response at low temperatures. We suggest that the observed anomalous electrodynamics may be related to quasiparticle traps or subgap states in the disordered TiN films. Finally, the electron density-of-states is derived from the pulse response. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.5723_arXiv.txt": {
        "abstract": "We describe the current status of solar neutrino measurements and of  the theory -- both neutrino physics and solar astrophysics -- employed in interpreting measurements.  Important recent developments include Super-Kamiokande's  determination of the $\\nu-e$ elastic scattering rate for $^8$B neutrinos to 3\\%; the latest SNO global analysis in which the inclusion of low-energy data from SNO I and II significantly narrowed the range of allowed values for the neutrino mixing angle $\\theta_{12}$; Borexino results for both the $^7$Be and pep neutrino fluxes, the first direct measurements constraining the rate of ppI and ppII burning in the Sun;  global reanalyses of solar neutrino data that take into account new reactor results on $\\theta_{13}$; a new decadal evaluation of the nuclear physics of the pp chain and CNO cycle defining best values and uncertainties in the nuclear microphysics input to solar models; recognition of an emerging discrepancy between two tests of solar metallicity, helioseismological mappings of the sound speed in the solar interior, and analyses of the metal photoabsorption lines based on our best current description of the Sun's photosphere; a new round of standard solar model calculations optimized to agree either with helioseismology or with the new photospheric analysis; and, motivated by the solar abundance problem, the development of nonstandard, accreting solar models, in order to investigate possible consequences of the metal segregation that occurred in the proto-solar disk.  We review this progress and describe how new experiments such as SNO+ could help us further exploit neutrinos as a unique probe of stellar interiors. ",
        "introduction": "\\label{sec:one} In 1958  \\cite{HJ58,HJ59} found that the cross section for ${}^3\\mathrm{He}+{}^4\\mathrm{He}\\rightarrow {}^7\\mathrm{Be}+\\gamma$ was about 1000 times larger than anticipated, so that in addition to the simplest ${}^3\\mathrm{He}+{}^3\\mathrm{He} \\rightarrow {}^4\\mathrm{He}+2\\mathrm{p}$ ppI termination of the pp chain (see Fig. \\ref{fig:cycles}), there might be significant branches to the ppII and ppIII cycles, and thus significant fluxes of ${}^7$Be and ${}^8$B solar neutrinos.  Despite the uncertainties that existed in 1958 -- the solar core temperature was poorly constrained by theory, and other nuclear physics important to the pp chain had not been resolved -- both \\cite{Cameron} and \\cite{WAF58} pointed out that it might therefore be possible to detect solar neutrinos using a radiochemical method Ray Davis had developed at Brookhaven \\citep{Davis55}.  While the endpoint of the main source of neutrinos from the ppI cycle, p+p$\\rightarrow$d+e$^+$+$\\nu_e$, is below the 811 keV threshold for $\\nu_e$+${}^{37}$Cl$\\rightarrow {}^{37}$Ar + e$^-$, most ${}^7$Be and $^8$B neutrinos are sufficiently energetic to drive this reaction.  In 1962 Fowler organized a team of young Caltech researchers -- John Bahcall, Icko Iben, and Dick Sears -- to begin the development of a solar model to more accurately predict the central temperature of the Sun and to estimate the rates of neutrino-producing reactions \\citep{BFIS}.   The history of these early developments is summarized in several sources \\citep{Account,Lande10,Haxton10}. By early 1964, following significant advances in the solar model and in the understanding of the nuclear physics of the pp chain and the ${}^{37}$Cl($\\nu_e$,e$^-$)${}^{37}$Ar reaction, \\cite{DavisPRL64} and \\cite{JNBPRL64} concluded that a measurement of solar neutrinos would be possible, were Davis to mount a detector 100 times larger than that he built at Brookhaven, in a site sufficiently deep to reduce backgrounds from high-energy cosmic ray muons to an acceptable level.  In April 1968 \\cite{DHH68} announced an upper bound on the solar neutrino capture rate for ${}^{37}$Cl of 3 SNU (1 SNU = 10$^{-36}$ captures/target atom/s), based on the initial running of a 100,000-gallon C$_2$Cl$_4$ detector that the collaborators had constructed on the 4850-ft level of the Homestake gold mine, in Lead, South Dakota.  This upper bound, nearly a factor of three below the rate predicted by the \\cite{BBS68} standard solar model (SSM), began a controversy that  took 30 years to resolve.  Twenty of these years passed without independent confirmation of the Davis result: as the Cl rate was a fraction of a count per day in 0.6 kilotons of organic liquid, other technologies with comparable sensitivity were not easily developed.  Because the Davis experiment was sensitive to a flux of neutrinos that varies sharply with the solar core temperature ($\\phi(^8$B) $\\sim T_C^{22}$ \\citep{Book}), the result could be accommodated by a variety of possible changes in the SSM having the net effect of reducing $T_C$ by $\\sim$ 5\\%.   But as additional constraints on solar neutrino fluxes were established by the Kamiokande \\citep{KamiokandeA}, SAGE \\citep{SAGE}, and GALLEX \\citep{GALLEX} collaborations, a more detailed pattern of fluxes emerged that was not easily reconciled with the SSM. In contrast, with the discovery of the MSW mechanism \\citep{MS1985,MS1986,Wolfa, Wolfb}, it became apparent that neutrino oscillations augmented by matter effects could account for the observations, even for a small mixing angle.  The conclusions of an Annual Reviews article from this period \\citep{Haxton95} captures the sense of excitement that with new experiments, a resolution of the solar neutrino problem was near.   In 1998 $\\nu_\\mu \\rightarrow \\nu_\\tau$ vacuum neutrino oscillations were discovered through the Super-Kamiokande collaboration's study of the zenith-angle dependence of atmospheric neutrino fluxes \\citep{SK98}.  While this result did not directly constrain the $\\nu_e$s produced by the Sun, the discovery was a game-changer, confirming a phenomenon originally suggested by \\cite{Pontecorvo67} as a possible explanation for the solar neutrino problem.  Finally, the Sudbury Neutrino Observatory (SNO) collaboration \\citep{SNO1,SNO2} measured both the $\\nu_e$ and heavy-flavor components of the solar neutrino flux arriving at Earth,  utilizing three different detection channels with varying sensitivities to charge and neutral currents.  The SNO collaboration measured the electron and heavy flavors components of the $^8$B solar neutrino flux, found that the total flux of neutrinos (summed over flavors) is in good agreement with the SSM prediction, and determined flavor mixing parameters that attributed the differential suppression of the pp, $^7$Be, and $^8$B fluxes deduced from previous experiments to the energy-dependent effects of matter on oscillations.  This review summarizes the basic physics of solar neutrinos, the work that led to the discoveries noted above, and the impact of recent and ongoing solar neutrino experiments on astrophysics and weak interactions, including \\begin{itemize} \\item Completion of phase III of the Super-Kamiokande experiment \\citep{SKIII} and preliminary results from Super-Kamiokande IV's low-threshold running \\citep{Smy}; \\item SNO's combined analysis of all three SNO phases \\citep{SNOCombined} and the collaboration's low-energy analysis of the data from SNO I and II \\citep{SNOLET}; \\item  Borexino's achievement of a 5\\% measurement of the $^7$Be flux, an initial result for the pep flux, and a limit on the CN neutrino contribution \\citep{Borexino,Borexinopep}; and \\item Daya Bay, Reno, and Double Chooz measurements of $\\theta_{13}$, impacting global analyses of solar neutrino data \\citep{DayaBay,Reno,DC}. \\end{itemize} In addition, a comprehensive survey of the nuclear physics of the pp chain and CNO cycle has been completed, yielding a new set of best values and uncertainties for the nuclear rates \\citep{SFII}.  The sound speed throughout most of the Sun's interior has been extracted from helioseismology to an accuracy $\\sim$ 0.1\\%, providing a stringent check on SSM predictions. More sophisticated 3D models of the solar atmosphere have been developed, significantly improving the agreement between predicted and observed absorption line-shapes and the consistency of abundance determinations from different atomic and molecular lines  \\citep{AGSS09} -- but also yielding a photospheric metal abundance $\\sim$ 30\\% below 1D values, leading to a conflict between SSMs employing the new abundances and solar parameters deduced from helioseismology. The SSM has been updated for the new nuclear reaction rates and alternative metallicities, and nonstandard models have been developed to explore accretion as a possible solution to the ``solar abundance problem\" \\citep{Serenelli09,Serenelli11}.  For three decades solar neutrino physics was defined by an incomplete knowledge of the neutrino fluxes and shortcomings in our understanding of neutrino flavor physics. We are now starting a new period, where precise spectroscopy of the full spectrum of solar neutrinos is possible, and where a clearer understanding of weak interactions has been obtained from a combination of astrophysical, reactor, and accelerator experiments. On one hand, this returns us to the roots of solar neutrino physics: with weak interaction uncertainties removed, solar neutrinos can be used to probe possible limitations in the SSM -- such as uncertainties in the Sun's initial composition and the absence of {\\it ab initio} treatments of mixing and other three-dimensional physics, including rotation and magnetic fields.  On the other hand, the neutrinos coming from the Sun remain important to fundamental physics: the spectral shapes and fluxes of the various sources are  known rather precisely, and low-energy neutrinos react with targets rather simply, giving us confidence that we can interpret measurements. Thus this review also considers the continuing role solar neutrinos could play in further constraining the Pontecorvo-Maki-Nakagawa-Sakata (PMNS) neutrino mass matrix and in probing matter effects and other environmental neutrino phenomena. ",
        "conclusions": ""
    },
    "1208/1208.6100_arXiv.txt": {
        "abstract": "{The gravitational magnification provided by massive galaxy clusters makes it possible to probe the physical conditions in distant galaxies that are of lower luminosity than those in blank fields and likely more representative of the bulk of the high-redshift galaxy population.} {We aim to  constrain the basic properties of molecular gas in a strongly magnified submm galaxy located behind the massive Bullet Cluster (1E~0657-56). This galaxy (SMM\\,J0658) is split into three images, with a total magnification factor of almost 100.}  {We used the Australia Telescope Compact Array (ATCA) to search for $^{12}$CO(1--0)~and $^{12}$CO(3--2)~line emission from SMM\\,J0658. We also used the SABOCA bolometer camera on the Atacama Pathfinder EXperiment (APEX) telescope to measure the continuum emission at 350~$\\mu \\mathrm{m}$.} {CO(1--0)~and CO(3--2)~are detected at {6.8}$\\sigma$ and {7.5}$\\sigma$ significance when the spectra toward the two brightest images of the galaxy are combined. From the CO(1--0)~luminosity we derive a mass of cold molecular gas of ${(1.8 \\pm 0.3) \\times 10^9} \\, M_{\\odot}$, using the CO to H$_2$ conversion factor commonly used for luminous infrared galaxies. This is $45 \\pm 25\\%$ of the stellar mass.  From the width of the CO lines we derive a dynamical mass within the CO-emitting region $L$ of ${(1.3 \\pm 0.4) \\times 10^{10}{ (L/1 {\\rm kpc}) }} \\, M_{\\odot}$. We refine the redshift determination of SMM\\,J0658~ to $z={2.7793 \\pm 0.0003}$. The CO(3--2)~to CO(1--0)~brightness temperature ratio is ${0.56_{-0.15}^{+0.21}}$, which is similar to the values found in other star-forming galaxies. Continuum emission at 350~$\\mu \\mathrm{m}$~from SMM\\,J0658~was detected with SABOCA at a signal-to-noise ratio of 3.6. The flux density is consistent with previous measurements at the same wavelength by the Herschel satellite and BLAST balloon-borne telescope. We study the spectral energy distribution of SMM\\,J0658~and derive a dust temperature of $33\\pm 5$~K and a dust mass of ${1.1_{-0.3}^{+0.8}\\times 10^{7} \\, M_{\\odot}}$.} {SMM\\,J0658~is one of the least massive submm galaxies discovered so far. As a likely representative of the bulk of the submm galaxy population, it is a prime target for future observations.}\\date{\\today} ",
        "introduction": "\\label{sec:introduction}  Molecular gas is the raw material from which stars form. Determining the amount of molecular material and the physical conditions of the molecular interstellar medium in distant galaxies is important to understand the cosmic history of star-formation and the evolution of galaxies (see reviews by \\citealt{SolomonVanden-Bout:2005aa} and \\citealt{Walter:2010aa}).  Since the first detection of carbon monoxide in a high-redshift galaxy twenty years ago (IRAS~F10214+4724 at $z=2.3$; \\citealt{BrownVanden-Bout:1991aa, SolomonRadford:1992aa}), CO and other molecules have been detected in different types of distant galaxies, indicating the presence of significant reservoirs of molecular gas: $> 10^8 M_\\odot$ in Lyman-break galaxies (e.g.,~\\citealt{StanwayBremer:2008aa}, \\citealt{RiechersCarilli:2010aa}), larger amounts (a few times $10^9 M_\\odot$) in some submillimeter-faint radio-selected starburst galaxies (e.g. \\citealt{ChapmanNeri:2008aa}), and even larger amounts (> $10^{10} M_\\odot$) in near-infrared--selected star-forming galaxies (e.g., \\citealt{DannerbauerDaddi:2009aa}), submillimeter galaxies (SMGs) (e.g., \\citealt{GreveBertoldi:2005aa,DaddiDannerbauer:2009aa,2011MNRAS.412.1913I}), and in quasars and quasi-stellar objects (e.g., \\citealt{AlloinKneib:2007aa,CoppinSwinbank:2008aa}). The SMG population contains extreme objects: high-redshift dust-obscured star-forming galaxies with rest-frame far-infrared luminosities larger that 10$^{12}L_\\odot$ and star-formation rate of about 1000 $M_\\odot$yr$^{-1}$ (e.g. \\citealt{SmailIvison:1997aa}, \\citealt{BlainSmail:2002aa}).  Whereas the brightest SMGs contribute largely to the star-formation rate in the Universe, they are not representative of the high-redshift SMG population as a whole: number counts indicate a steep increase of the SMG population with decreasing flux density, $S_\\nu$: the space density of SMGs is about $10^4$ deg$^{-2}$ for sources with $S_\\nu > 1$~mJy (e.g. \\citealt{HughesSerjeant:1998aa,ChapmanRichards:2001aa,SmailIvison:2002aa}), and it is about ten times larger for galaxies with $S_\\nu> 0.1$~mJy (e.g. \\citealt{Knudsenvan-der-Werf:2008aa}). The abundant sub-mJy population is difficult to detect { in observations of relatively poor angular resolution ($15\\arcsec-30\\arcsec$) because of confusion noise}.  Most detections so far were possible because of gravitational lensing by a foreground galaxy or a galaxy cluster, that not only brightens a source but also provides an effective increase in angular resolution, which lowers the confusion limit (e.g. \\citealt{SmailIvison:1997aa,SmailIvison:2002aa,ChapmanSmail:2002aa,Knudsenvan-der-Werf:2005aa,KnudsenBarnard:2006aa,Knudsenvan-der-Werf:2008aa,JohanssonHorellou:2010ab,WardlowSmail:2010aa, RexAde:2009aa,EgamiRex:2010aa,JohanssonSigurdarson:2011aa}).  With the Atacama Large Millimeter/Submillimeter Array (ALMA), the Large Millimeter Telescope (LMT) and the Cerro Chajnantor Atacama Telescope (CCAT), the situation will be transformed, making it possible to observe faint sources directly. Until then, observations of highly magnified sources provide a first glimpse into the bulk of the SMG population.  Few sub-mJy SMGs have been discovered so far: there are seven in \\citet{Knudsenvan-der-Werf:2008aa}'s sample, and five in \\citet{CowieBarger:2002aa}'s sample. \\citet{SmailIvison:2002aa} estimated number counts down to 0.25~mJy, but the low-end counts were inferred from a Monte Carlo analysis and lower limits on the magnification of several faint sources. To allow { detection} in a reasonable amount of observing time, the submm flux density (before correcting for magnification) should preferrably be larger than 10~mJy, implying a magnification higher than 10. One SMG that satisfies those criteria is the galaxy at $z\\sim 2.5$ lensed by the massive cluster Abell~2218 (\\citealt{Kneibvan-der-Werf:2004aa}, \\citealt{KneibNeri:2005aa}). This source has a total magnification of 45 and an intrinsic submm flux of 0.8~mJy.  CO(3--2) was detected in all three images and CO(7--6) in the brightest one, and a gas mass of $4.5\\times 10^9$ and a star formation rate of about $500 \\, M_{\\odot} \\, \\mathrm{yr}^{-1}$ were inferred \\citep{KneibNeri:2005aa}.  Another sub-mJy SMG is the recently discovered galaxy SMM~J065837.6--555705 (hereafter SMM\\,J0658) at $z=2.79$ situated {near the caustic line of the} Bullet Cluster (1E~0657-56) at $z\\sim 0.3$, {and magnified about 100 times} (\\citealt{Bradac:2006fj,2008MNRAS.390.1061W,GonzalezClowe:2009aa}, \\citealt{GonzalezPapovich:2010aa} (hereafter G10), \\citealt{JohanssonHorellou:2010ab}). It is the subject of this paper. The galaxy classifies as a Luminous Infrared Galaxy (LIRG), with an intrinsic far-infrared luminosity of a few times 10$^{11} L_\\odot$. The source was also detected by BLAST and Herschel; its infrared spectral energy distribution is consistent with that of a dusty starburst galaxy \\citep{RexAde:2009aa,RexRawle:2010aa}. It is the brightest source in our APEX LABOCA 870~$\\mu \\mathrm{m}$~survey of gravitationally lensed submm galaxies \\citep{JohanssonSigurdarson:2011aa}, in which it was detected with a significance of $\\sim 30\\sigma$. When corrected for the gravitational magnification, it appears to be one of the intrinsically faintest submm galaxies detected so far ($S_{\\mathrm{870 \\, {\\mu}m}} \\sim 0.5$~mJy). The galaxy is strongly lensed, and three images, A through C, were identified in infrared Spitzer images (\\citealt{GonzalezClowe:2009aa}, G10). Images A and B are separated by 8\\arcsec~and have individual magnification factors of $\\sim30$ and $\\sim70$ (G10). The third image, C, lies 30\\arcsec~away from the centroid of images A and B, and its magnification factor is smaller ($< 10$), making it too faint for detection in the submm and mm observations discussed here. Recently, a faint arc extending between images A and B was discovered in Hubble Space Telescope near-infrared images (G10). G10 also presented the first spectroscopic redshift of SMM\\,J0658, $z=2.791\\pm 0.007$, derived from Polycyclic Aromatic Carbon (PAH) bands. They also reported the detection of two rotational lines of H$_2$, and derived a warm molecular gas mass of $2.2_{-0.8}^{+17} \\times 10^8 \\; (\\mu_{\\mathrm{AB}} / 100 )^{-1} \\, M_{\\odot}$, where $\\mu_{\\mathrm{AB}}$ is the total magnification, and a gas temperature of $377_{-85}^{+68}$ K.  In this paper, we present the first detections of CO(1--0) and CO(3--2) in SMM\\,J0658. The observations were done with the Australia Telescope Compact Array (ATCA). We use the CO detections to {constrain the basic properties of the cold} molecular gas in SMM\\,J0658. To complement the existing Herschel observations we also present APEX SABOCA observations of the 350~$\\mu \\mathrm{m}$~continuum that we use to quantify the dust properties.  Throughout the paper, we adopt the following cosmological parameters: a Hubble constant $H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, a matter density parameter $\\Omega_0 = 0.27$, and a dark energy density parameter $\\Omega_{\\Lambda0} = 0.73$. In this cosmology, $z=2.8$ corresponds to an angular-diameter distance of 1650~Mpc, a luminosity distance of 23800~Mpc and a scale of 8.0~kpc/arcsec\\footnote{We used Ned Wright's cosmology calculator \\citep{Wright:2006aa} available at \\texttt{http://www.astro.ucla.edu/{\\textasciitilde}wright/cosmocalc.html}.}. ",
        "conclusions": "\\label{sec:discussion}  This paper presents a study of the molecular gas and dust in one of the least massive high-redshift galaxies observed so far, the $z=2.8$ submm galaxy SMM\\,J0658~lensed by the Bullet Cluster.  \\begin{itemize} \\item We detected for the first time rotational transitions of CO from SMM\\,J0658. The CO(1--0)~to CO(3--2)~brightness temperature ratio is comparable to that observed in other high-redshift star-forming galaxies. \\item We revise the redshift estimated by \\citet{GonzalezPapovich:2010aa} from $z=2.791\\pm 0.007$ to $z=2.7793\\pm 0.0003$. \\item The mass of cold molecular gas is estimated to be between $1/3$ and $2/3$ of the total baryonic mass. \\item From the linewidths of the CO(3--2)~transition we derive a dynamical mass of $(1.3 \\pm 0.4)\\times 10^{10} \\, M_{\\odot}$, {for a CO-emitting disk with a physical size of 1 kpc}. \\item The derived molecular and dynamical masses are consistent with a galaxy less massive than the Milky Way. SMM\\,J0658~is also less massive than most star-forming galaxies observed at high redshift so far (e.g. \\citealt{TacconiGenzel:2010ab,GreveBertoldi:2005aa,SwinbankSmail:2010aa}). \\item Using ground-based 350~$\\mu \\mathrm{m}$~imaging we detected continuum radiation from SMM\\,J0658. The signal-to-noise ratio is only 3.6, but {the detection} makes it possible to confirm that the emission previously seen in the larger beams of BLAST and Herschel (\\citealt{RexAde:2009aa}, \\citealt{RexRawle:2010aa}) comes from SMM\\,J0658.  The value of the flux density measured by SABOCA is consistent with the Herschel and BLAST values. \\item Owing to the extremely high magnification, this galaxy is a target of choice for further studies of the properties of intrinsically faint high-redshift galaxies, for example with ALMA.  \\end{itemize}"
    },
    "1208/1208.3455_arXiv.txt": {
        "abstract": "In this paper we present the results of the radio light curve and X-ray observations of broad-lined Type Ic SN 2007bg. The light curve shows three distinct phases of spectral and temporal evolution, implying that the SNe shock likely encountered at least 3 different circumstellar medium regimes. We interpret this as the progenitor of SN 2007bg having at least two distinct mass-loss episodes (i.e., phases 1 and 3) during its final stages of evolution, yielding a highly-stratified circumstellar medium. Modelling the phase 1 light curve as a freely-expanding, synchrotron-emitting shell, self-absorbed by its own radiating electrons, requires a progenitor mass-loss rate of $\\dot{M}\\approx1.9\\times10^{-6}(v_{w}/1000\\;\\rm{km\\,s}^{-1})$ \\Msunpyr\\ for the last $t\\sim20(v_{w}/1000\\;\\rm{km\\,s}^{-1})$ yr before explosion, and a total energy of the radio emitting ejecta of $E\\approx1\\times10^{48}$ erg after $10$ days from explosion. This places SN 2007bg among the most energetic Type Ib/c events. We interpret the second phase as a sparser \"gap\" region between the two winds stages. Phase 3 shows a second absorption turn-on before rising to a peak luminosity $2.6$ times higher than in phase 1. Assuming this luminosity jump is due to a circumstellar medium density enhancement from a faster previous mass-loss episode, we estimate that the phase 3 mass-loss rate could be as high as $\\dot{M}\\la4.3\\times10^{-4}(v_{w}/1000\\;\\rm{km\\,s}^{-1})$ \\Msunpyr. The phase 3 wind would have transitioned directly into the phase 1 wind for a wind speed difference of $\\approx2$. In summary, the radio light curve provides robust evidence for dramatic global changes in at least some Ic-BL progenitors just prior ($\\sim 10-1000$ yr) to explosion. The observed luminosity of this SN is the highest observed for a non-gamma-ray-burst broad-lined Type Ic SN, reaching $L_{8.46\\,\\rm{GHz}}\\approx1\\times10^{29}$ \\ergphzps, $\\sim567$ days after explosion. ",
        "introduction": "Core collapse supernovae (SNe) mark the spectacular death of massive stars \\citep*{1986ARA&A..24..205W}. Among these SNe Types Ib and Ic (collectively referred to as Ib/c due to their perceived similarities) are associated with the explosion of massive stars which were stripped of their H envelopes by strong stellar winds. Types Ib show He features in their spectra, whereas Types Ic do not. Our current understanding of Type Ib/c SNe suggests that Wolf-Rayet (WR) stars are a primary progenitor path \\citep*{2003PASP..115....1V}, although the relative numbers of WR stars fall well short of the observed Type Ib/c rates and thus a lower-mass binary progenitor path is most likely required to account for a significant fraction of these events \\citep[e.g.,][]{1992ApJ...391..246P,2004MNRAS.349.1093R,2011ApJ...740...79W,Smith2011}. Archival imaging of pre-SNe locations, however, have failed to produce direct associations between Type Ib/c SNe and their possible progenitors \\citep[e.g.,][]{2009ARA&A..47...63S}. Thus we must rely both on theory and indirect information obtained after the explosion. Fortunately, X-ray and radio emission are detectable by-products of the shock interaction between the high-velocity ejecta and the low-velocity progenitor wind in SNe (whereas the optical emission is generated by lower-velocity ejecta and radioactive decay). These observables crudely scale with the density of the circumstellar medium (CSM), and hence provide constraints on the wind properties of the SN progenitor \\citep[e.g.,][]{1982ApJ...258..790C,1982ApJ...259..302C}. Because massive stars are expected to have strong episodic mass-loss \\citep[e.g.,][]{1994PASP..106.1025H} and line driven winds \\citep*[e.g.,][]{Vink2000}, radio monitoring of Type Ib/c SNe provide interesting constraints on possible progenitor types. The past decade has witnessed increased attention on Type Ib/c SNe due to the association of Type Ic SN 1998bw with gamma-ray burst (GRB) 980425 \\citep{1998Natur.395..670G}, as these objects may help elucidate how the central engines of both GRBs and their associated SNe work. Following this historic connection, several other Type Ic SNe have been linked to GRBs \\citep[e.g.,][]{Hjorth2003,Kawabata2003,Matheson2003,2003ApJ...591L..17S,2005ApJ...619..994F, 2006ApJ...645L..21M,2006ARA&A..44..507W,Bufano2012};  all of these objects show broad emission lines in their optical spectra, with velocities above $15000$ \\kmps. To date so far 5 GRBs have been spectroscopically confirmed to have a SNe counterpart, while 14 show photometric evidence of a SNe counterpart \\citep*{Hjorth2011}, and hence a larger sample is needed in order to shed some light on the explosion mechanisms involved. Given the large number of SNe Type Ic compared to GRBs, one key question is whether these SNe could simply be off-axis GRBs. Initial efforts concluded that the vast majority were not \\citep[e.g.,][]{2006ApJ...638..930S}. Recently new evidence was presented for the detection of at least one possible off-axis GRB with no direct observation of the gamma-ray emission: SN 2009bb \\citep{2010Natur.463..513S}; as well as the more controversial SN 2007gr \\citep{2010Natur.463..516P,2010ApJ...725..922S,2011ApJ...735....3X}.  Here we report on a particularly interesting broad-lined Type Ic, SN 2007bg, which was optically discovered on 2007 April 16.15 UT at $\\alpha=11^{\\rm h} 49^{\\rm m} 26.18^{\\rm s}$, $\\delta=+51^{\\circ} 49^{\\prime} 21.8^{\\prime\\prime}$ (J2000) at redshift $z=0.0346$ \\citep*{2007CBET..927....1Q} ($d=152$ Mpc, $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm{M}}=0.3$, and $\\Omega_{\\rm{\\Lambda}}=0.7$ in a $\\Lambda$CDM cosmology). A spectrum taken on 2007 April 18.3 with the 9.2 m Hobby-Everly Telescope let \\citet{2007CBET..927....1Q} classify SN 2007bg as a Type Ic broad-lined (Ic-BL) supernova. The classification of this SN as a Type Ic and the fact that it resides in a faint host encouraged \\citet*{2008ApJ...673..999P} to think about this object as a likely off-axis GRB. \\citet*{2009ATel.2065....1P} also noted that SN 2007bg was one of the brightest radio SNe one year after explosion, making it an even better candidate for an off-axis GRB. However, based on the radio light curve of SN 2007bg known at the time, \\citet{2009ATel.2066....1S} contended that the observed radio emission could be explained by the presence of a density enhancement in the CSM, rather than a GRB like central engine driving the explosion. Using the archival radio data for SN 2007bg, we further address this issue below.  The article outline is the following. In \\S\\ref{s:obs} we introduce the observations taken with the Very Large Array (VLA). Our results and modelling of the radio light curves are presented in \\S\\ref{s:results}. Finally in \\S\\ref{s:discussion} we summarise our findings and provide a brief discussion on the future of radio observations of SNe. ",
        "conclusions": "\\label{s:discussion}  In this article we presented the radio light curves and X-ray observations of Type Ic-BL SN 2007bg. The radio emission is characterised by three different phases. Phase 1 being characterised by a SSA turn-on. Based on the brightness temperature of SN 2007bg the expansion speed of the radio emitting region is $v=0.19\\pm0.02c$. An expansion speed similar to what has been found in the small sample of radio-emitting Type Ib/c SNe.  The derived mass-loss rate from the progenitor during phase 1 is $\\dot{M}\\approx1.9\\times10^{-6}(v_{w}/1000\\;\\rm{km\\,s}^{-1})$ \\Msunpyr, a low mass-loss rate among the observed in WR stars. During phase 2 we observe a drop in the flux density while the spectral index seems to follow a smooth evolution. This drop is presumably due to the radio ejecta entering a region of lower CSM density, probably arising because of a change in the mass-loss rate of the progenitor or a variable wind speed. Then during phase 3 we observe a rise in the radio flux density along with a re-absorption of the radio emission. The shallow observed slope of the re-absorption could be explained by small clumps of material in the progenitor wind which would become more strongly self-absorbed than the surrounding CSM when shocked and produce a shallower optically-thick spectrum. From our modelling of the radio light curves during phase 3 we put an upper limit on the mass-loss rate of $\\dot{M}\\la4.3\\times10^{-4}(v_{w}/1000\\;\\rm{km\\,s}^{-1})$ \\Msunpyr. This second wind component could arise from a different stellar evolution phase of the progenitor prior to explosion, possibly a LBV phase. Or it could be the effect of stellar rotation on the stellar wind properties. If the radio flux density from SN 2007bg declines in time like $f_{\\nu}\\propto t^{-1}$. It should still be detectable, with a flux density at $22.5$ GHz of $\\sim500$ \\mJy, and higher at lower frequencies. If the progenitor star went through more mass-loss episodes, then the observed flux density would deviate from this value. Recent observations taken with the Karl G. Jansky Very Large Array (JVLA) should help to establish this. Very few broad-lined type Ic SNe have radio detections, and none of them show signs of a complex CSM like that of SN 2007bg.  These different wind components, along with the derived mass-loss rates will eventually help to constrain the progenitors of Type Ib/c SNe, from their inferred pre-SNe properties. Currently two main evolutionary paths has been proposed to account for the rate of these explosions: a single WR origin, and a binary system. However, the single WR scenario itself still needs significant clarification. For instance, what effects do evolution (mass-loss, rotation, etc.) and metallicity have in shaping the CSM around massive type Ib/c progenitors and ultimately their explosions. And more generally, how does this fold into the connection between type II's, type Ib/c, and GRBs.  Compounded by their relative rarity, a key impediment to improving our knowledge of broad-lined Ic SNe has been the lack of high-quality, multi-frequency observations from early through late epochs. Sensitive high-frequency radio observations of Ic SNe within 1-2 days of explosion are needed to efficiently identify the best objects for follow-up and allow further refinement of the physical properties surrounding these unique explosions. Both ALMA and the newly retooled JVLA will play critical roles here, providing rigorous probes of the synchrotron-emitting shock (to constrain evolution of the shock velocity, any potential beaming, etc.) and ultimately new insights into the properties of their CSM (such as density, variability, clumpiness, and perhaps even dust content for these massive systems). Sensitive, long-term X-ray and optical/NIR spectroscopic follow-up are also needed in order to break potential degeneracies and provide consistency checks against the shock-CSM interaction models. If such complete datasets can be assembled in the next several years, they may afford us a leap in our physical understanding of SNe and GRBs."
    },
    "1208/1208.2179_arXiv.txt": {
        "abstract": "We report on new simulations of the transport of energetic protons originating from the decay of energetic neutrons produced in solar flares.  Because the neutrons are fast-moving but insensitive to the solar wind magnetic field, the decay protons are produced over a wide region of space, and they should be detectable by current instruments over a broad range of longitudes for many hours after a sufficiently large gamma-ray flare.  Spacecraft closer to the Sun are expected to see orders-of-magnitude higher intensities than those at the Earth-Sun distance. The current solar cycle should present an excellent opportunity to observe neutron-decay protons with multiple spacecraft over different heliographic longitudes and distances from the Sun. ",
        "introduction": "Solar flares remain some of the best laboratories in the solar system for studying particle acceleration, production and interaction.  The particles accelerated in flares either interact with the solar chromosphere, producing observable X-ray and $\\gamma$-ray emission, or escape into the interplanetary medium.  In 1951, Biermann, Haxel and Schluter first suggested that solar flares could produce energetic neutrons, and evidence of these neutrons was subsequently observed in the 2.2-MeV neutron-capture $\\gamma$-ray line, first by \\inlinecite{chu73} and many times since.  The first direct observations of solar-flare neutrons were made by \\inlinecite{chu82}.  Detailed models of neutron production, transport, interaction and decay in flares have been developed ({\\it e.g.} \\opencite{lin65}, \\opencite{hua87}, \\opencite{mur07}), allowing the number of interacting protons and escaping neutrons from a particular flare to be deduced from the 2.2 MeV line fluence.  The neutrons which escape the Sun undergo beta decay with a mean lifetime at rest of 886 seconds (about 14 minutes):  \\begin{equation} n^{0}~\\longrightarrow~p^{+}+e^{-}+\\overline{\\nu}_{e} \\end{equation} creating neutron-decay protons and electrons in the heliosphere.  The escaping neutrons and their decay products have been observed by both ground-based detectors and space-based instruments near the Earth and around the heliosphere (\\opencite{chu82}, \\opencite{deb83}).  Since the travel time between the Sun and Earth is $\\geq$20 minutes for particles with energies less than 100 MeV, most of the neutrons decay at some time during transit to 1 AU.  The neutron-decay protons \\cite{eve83a, eve90} and also electrons \\cite{dro96} typically arrive before the original flare-accelerated protons (protons of the same energy accelerated in the flare site rather than originating from neutron decay), and before particles accelerated by the shock driven by the associated coronal mass ejection (CME), because they make part of their journey as neutrons that travel in a straight line and also do not lose energy.  Following the early work by Roelof (1966), several studies have investigated the effects  that transport in the interplanetary magnetic field (IMF) can have on observations of neutron-decay protons, including their time-intensity profiles, energy spectra, and pitch-angle distributions (PADs). The first measurements of neutron-decay protons found the PADs were nearly isotropic, and, by fitting various aspects of the data, obtained scattering mean free-paths of $\\lambda$ = 0.3--0.5 AU \\cite{eve83a}, $\\lambda$ = 0.2 $\\pm$ 0.005 AU \\cite{eve83b}, and $\\lambda\\sim$ 0.3 AU \\cite{eve85}.  In a more complete treatment \\inlinecite{ruf91} modeled the transport of neutron-decay protons with 25 to 150 MeV and fit the data from events summarized by \\inlinecite{eve85}.  He considered pitch-angle scattering and adiabatic focusing of the protons and, by fitting ``distance-traveled'' distributions deduced from data obtained by the {\\it Third International Sun-Earth Explorer} (ISEE-3), obtained the best-fit scattering mean free-paths for two events of $\\lambda$ = 0.26 $\\pm$ 0.05 AU and $\\lambda$ = 0.37 $\\pm$ 0.06 AU.  Recently, \\inlinecite{agu11} modeled the transport of neutrons and their decay products to compare with observations of Solar Cycle 23 events.  Previous {\\it in situ} studies of neutrons and their decay products were limited to a single vantage point, but new multi-point instrumentation is providing a different perspective on the geometry of energetic particles in the solar system.  The canonical view \\cite{rea91,nit06} holds that CME shocks distribute particles widely in longitude, and flare-associated protons and electrons cover a narrow ($\\sim$40$^\\circ$) range in longitude.  However, recent studies have found that small flares can produce particles over $\\sim$80$^\\circ$ latitude \\cite{wib06, wie09}, while new imaging technology is routinely showing previously-undetectable CMEs associated with small flares. Understanding the longitude-distribution of neutron-decay protons may help unravel the diverse sources of energetic particles in the heliosphere.  The twin spacecraft of the {\\it Solar TErrestrial RElations Observatory} (STEREO, \\opencite{kai08}) were launched in 2006 into orbits that lead (STEREO-Ahead) and trail (STEREO-Behind) the Earth. The {\\it Advanced Composition Explorer} (ACE, \\opencite{sto98}), {\\it Solar and Heliospheric Observatory} (SOHO, \\opencite{dom83}), and {\\it Wind} spacecraft \\cite{acu95} are in orbit around Earth's L1 Lagrangian point, making multi-point studies covering more than 180$^\\circ$ in longitude possible in the current solar cycle. ",
        "conclusions": "The current solar cycle should present an excellent opportunity to observe neu\\-tron-decay protons with multiple spacecraft over different heliographic longitudes and distances from the Sun.  Our simulation suggests that these neutron-decay protons are very rapidly distributed over more than 180$^\\circ$ in longitude and should be detectable by current instrumentation at more than one longitude following large $\\gamma$-ray events, provided the pre-event energetic particle background is sufficiently quiet. Spacecraft closer to the Sun such as {\\it MESSENGER}, and eventually {\\it Solar Probe Plus} and {\\it Solar Orbiter}, will observe even higher intensities of neutron-decay protons.  While the decay protons and electrons with energies greater than a few MeV quickly escape into the heliosphere, lower-energy protons linger close to the Sun, where they are available to be accelerated by a CME shock into high-energy particles.  \\begin{acks} This work was supported by NASA under grants NNX08AI11G and NNX11AO75G, and subcontract SA2715-26309 from UC Berkeley under NASA contract \\linebreak NAS5-03131. We appreciate the assistance of R. A. Leske and C. M. S. Cohen, and discussions with R. Murphy. We also thank the anonymous reviewer for a number of helpful comments and suggestions. \\end{acks}"
    },
    "1208/1208.6078.txt": {
        "abstract": "We present extensive optical observations of a Type IIn supernova (SN) 2010jl for the first 1.5 years after the discovery. The $UBVRI$ light curves demonstrated an interesting two-stage evolution during the nebular phase, which almost flatten out after about 90 days from the optical maximum. SN 2010jl has one of the highest intrinsic H$\\alpha$ luminosity ever recorded for a SN IIn, especially at late phase, suggesting a strong interaction of SN ejecta with the dense circumstellar material (CSM) ejected by the progenitor. This is also indicated by the remarkably strong Balmer lines persisting in the optical spectra. One interesting spectral evolution about SN 2010jl is the appearance of asymmetry of the Balmer lines. These lines can be well decomposed into a narrow component and an intermediate-width component. The intermediate-width component showed a steady increase in both strength and blueshift with time until t $\\sim$ 400 days after maximum, but it became less blueshifted at t $\\sim$ 500 days when the line profile appeared relatively symmetric again. Owing to that a pure reddening effect will lead to a sudden decline of the light curves and a progressive blueshift of the spectral lines, we therefore propose that the asymmetric profiles of H lines seen in SN 2010jl is unlikely due to the extinction by newly formed dust inside the ejecta, contrary to the explanation by some early studies. Based on a simple CSM-interaction model, we speculate that the progenitor of SN 2010jl may suffer a gigantic mass loss ($\\sim$ 30-50 M$_{\\odot}$) in a few decades before explosion. Considering a slow moving stellar wind (e.g., $\\sim$ 28 km s$^{-1}$) inferred for the preexisting, dense CSM shell \\citep{smith11b, and11} and the extremely high mass-loss rate (1-2 M$_{\\odot}$ yr$^{-1}$), we suggest that the progenitor of SN 2010jl might have experienced a red supergiant stage and explode finally as a post-red supergiant star with an initial mass above 30-40 M$_{\\odot}$. ",
        "introduction": "Type II supernovae (SNe II) have long been thought to arise from the death of massive stars with $M >$ 8~M$_{\\odot}$ when the nuclear burning does not provide thermal pressure to support the star (e.g., \\citet{smar09}). Observationally, the presence of hydrogen in the spectra of SNe II are the distinguishing signatures of this class (see \\citet{fili97} for a review of SN classification), and they are usually fainter than the type Ia events. Type IIn supernovae (SNe IIn) represent a distinct subclass of SNe II that are characteristic of prominent narrow emission lines of hydrogen \\citep{sch90, fili91a, fili91b}. They can also have intermediate-width and/or broader components in addition to the narrow component. The narrow component usually has a full-width-half-maximum (FWHM) velocity less than a few hundred km s$^{-1}$ and is believed to arise from the pre-shocked, photo-ionized circumstellar material (CSM) shell around the exploding star \\citep{cf94}; while the intermediate- or broad-width component has an FWHM velocity of a few thousand km s$^{-1}$, and is formed from the dense post-shocked CSM shell. The presence of high-density CSM shells around some SNe IIn is favored by the detections of strong X-ray and radio emission after their explosion.  Type IIn supernovae are the brightest subgroup of all the SNe II, but they are rare events, accounting for $<$10\\% of all the SNe II \\citep{smar09, lwd11, smith11a}. The well-observed, representative SNe IIn sample are SNe 1988Z \\citep{tur93}, 1994W \\citep{soll98, chu04}, 1998S \\citep{fas00, liu00, pool02}, 2006tf \\citep{smith08}, and SN 2006gy \\citep{smith07a}, which exhibit diverse observational properties. For example, the maximum-light absolute magnitude can vary from $\\sim$ $-$19.0 mag (for SNe 1988Z and 1994W) to $\\sim$ $-$22.0 mag (for SN 2006gy) in the broadband $R$. The strength and profile of the emission lines such as the H$\\alpha$ lines also differ significantly among these SNe IIn \\citep{smith08}. This indicates a large spread for the properties of the progenitor stars for this subgroup of core-collapse SNe (see a summary by \\citet{kie12}).  The dense circumstellar medium can be formed due to a high mass-loss rate of massive stars \\citep{chu04, gal09}. The luminous blue variable (LBV) stars are very bright, blue, hypergiant variable stars that experiences frequent eruptions \\citep{hd94}, and has been proposed to be the possible progenitors of some type IIn SNe. One well-known case is the luminous type IIn SN 2005gl. Pre-explosion $Hubble$ $Space$ $Telescope$ (HST) images show that the bright point source at the position of SN 2005gl is likely an LBV-like star, and the speed of the pre-shocked progenitor wind is also consistent with an LBV star \\citep{gal07, gal09}. Another case is SN 2006tf which was also suggested to have a pre-shocked wind speed of a LBV-like star, but not of red supergiants or Wolf-Rayet stars \\citep{smith08}. On the other hand, it is also proposed that the WR star with fast winds lasting for a few thousand years might be the probable progenitors of some SNe IIn \\citep{dvv10, dvv11}. Note that not all of the type IIn supernovae show strong radio or X-ray emission suggests that not all of them are undergoing a strong interaction with a dense CSM. This suggests that SNe IIn may have more than one type of progenitor.  SN 2010jl was discovered on November 02.06 UT \\citep{np10} in an irregular galaxy UGC 5189A, and is one of the brightest supernova recorded in 2010 (13.5 mag on 2010 November 3). The early observation shows that it is a type IIn in the relatively young phase, with distinguished signatures of multiple hydrogen emission lines in the spectrum \\citep{ben10}. At the optical position of the SN, an X-ray point-source was also detected on November 5 with a 6.5-$\\sigma$ significance of source detection \\citep{imm10}. Using the archival WFPC2 images of $Hubble$ $Space$ $Telescope$ taken at about 10 years prior to explosion, \\citet{smith11b} suggests that the progenitor of SN 2010jl be consistent with a massive star with an initial mass above 30~M$_{\\odot}$ but the exact nature of the progenitor is not conclusive. The optical linear spectropolarimetry of SN 2010jl obtained at two weeks after the discovery shows a continuum polarization at a level of 1.7-2.0\\%, indicating that the explosion has a substantial asphericity \\citep{pat11}. Near infrared observations were obtained at about 90 days after the explosion with the $Wisconsin$ $Indiana$ $Yale$ $NAOA$ telescope for the $JHK_{s}$ bands and the $Spitzer$ IRAC for two mid-infrared bands (with the wavelength centering at 3.6 and 4.5$\\mu$m, respectively) \\citep{and11}. These observations revealed a significant infrared (IR) excess for SN 2010jl, suggesting that a large amount of dust may exist around this supernova before explosion.  In this paper, we present extensive observations of the SN IIn 2010jl in optical bands, providing another well-observed example with which to compare other SNe IIn. \\citet{smith12} and \\citet{sto11} have studied some of the optical properties of SN 2010jl, but our excellent and independent dataset of both photometry and spectroscopy allow us to provide better constraints on the properties of SN 2010jl. Our observations and data reduction are described in $\\S$ 2, while $\\S$ 3 presents the light curves and spectra, and the spectral evolution (especially the hydrogen emission lines) is given in $\\S$ 4. In \\S 5 we present the discussions about the progenitor of SN 2010jl. Our summaries and conclusions are given in $\\S$ 5. ",
        "conclusions": "This paper presents the photometric and spectroscopic observations of the luminous type IIn Supernova 2010jl. Our data were mainly collected with the 0.8-m TNT and the 2.16-m telescope locating at the Xinglong Observatory of NAOC, spanning a period of about 1.5 years after the explosion. The main results inferred from our data and analysis for SN 2010jl and its progenitors are summarized below.  1. The $UBVRI$ light curves underwent a two-stage evolution, with a break at t $\\sim$ 90 days from the maximum. Before t $\\sim$ 90 days, the decay rates are in a range of about 0.7 $-$ 1.2 mag (100 days)$^{-1}$, with steeper slope in the bluer bands; while they become much smaller after t $\\sim$ 90 days from maximum due to that the light curves flatten out during this period, with values of 0.1-0.3 mag (100 days)$^{-1}$. Such a notable change in the decay rates likely relates to a significant contribution from CSM-interactions at the late phases, as also indicated by the strong emission of H$\\alpha$ lines seen in the late-time spectra of SN 2010jl. We also constructed the pseudo-bolometric light curve with our multicolor photometry. The light curve at early phases (t $<$ 90 days) can be well fit by a hypothetical $^{56}$Co decay of 3.4 M$_{\\odot}$, while the late-time evolution complies with an ejecta-CSM interaction model. This perhaps indicates that radioactive decay powers SN 2010jl at early phase and CSM interaction provides the energy at later phase.  2. We have also studied some of the important spectroscopic characteristics of SN 2010jl such as H$\\alpha$, H$\\beta$, and He I lines. These lines, especially the H lines appear very strong in the spectra of SN 2010jl. For example, the H$\\alpha$ emission is estimated to have a luminosity of about 0.6-1.0$\\times$10$^{42}$ erg s$^{-1}$ before t $\\sim$ 400 days. This is apparently stronger than that of other type IIn in the literature, especially at late phase. Another interesting feature about the line profiles of H$\\alpha$ and H$\\beta$ is that both can be well decomposed into a narrow component and intermediate-width component. The narrow component is likely produced by the outer, pre-shock CSM shell, with an FWHM velocity of about 700 km s$^{-1}$; the intermediate-width component is perhaps due to the post-shock CSM shells, with an FWHM velocity of $\\sim$ 2000-3400 km s$^{-1}$, depending on the phases. The intermediate-width component showed a progressively increased strength and blueshift with time until t $\\sim$ 400 days. However, an decrease in the strength and the blueshift of the H$\\alpha$ and H$\\beta$ lines was also noticed in the t $\\sim$ 515 day spectra. These changes, together with the flat evolution of the late-time light curves, suggest that the systematic blueshift of the intermediate-width component of the H lines may not be due to the newly formed dust in the post-shock region. Instead, the above changes of the line profiles could be a natural result of CSM-interactions: the intermediate-width component of H lines will appear stronger and more blueshifted with the acceleration of the CSM shells via the successive collision from the ejecta at early times, while the strength and blueshift of this component would decrease when the CSM-interactions became weak at late times. Nevertheless, a detailed quantitative analysis is needed for a test of such a hypothesis.  3. With our extensive light curves and spectral data, we further derived the physical parameters such as the mass-loss rate and stellar-wind flow timescale which allow us to place better constraints on the nature of the progenitor of SN 2010jl. The high mass-loss rate (e.g. $\\sim$ 1 M$_{\\odot}$) and a short flow timescale of the stellar wind (e.g., $\\sim$ 40 years), obtained for the progenitor of SN 2010jl, indicate that the progenitor star must undergo a tremendous amount of mass loss within a few decades before explosion. This makes us believe that SN 2010jl may have a massive progenitor like LBVs. However, the fact that the pre-shock CSM has a slow expansion velocity (e.g., $\\sim$ 28 km s$^{-1}$) comparable to that of some red supergiants suggests that SN 2010jl may explode at a post-red supergiant stage. The post-red supergiant can correspond to a group of LBVs with lower luminosity and a mass range of 30-60 M$_{\\odot}$. Thus a post-red supergiant progenitor might be possible for SN 2010jl."
    },
    "1208/1208.0894_arXiv.txt": {
        "abstract": "{ The most natural region of cosmologically compatible dark matter relic density in terms of low fine-tuning in a minimal supersymmetric standard model with nonuniversal gaugino masses is the so called bulk annihilation region. We study this region in a simple and predictive SUSY-GUT model of nonuniversal gaugino masses, where the latter transform as a combination of singlet plus a nonsinglet representation of the GUT group SU(5). The model prediction for the direct dark matter detection rates is well below the present CDMS and XENON100 limits, but within the reach of a future 1Ton XENON experiment. The most interesting and robust model prediction is an indirect detection signal of hard positron events, which resembles closely the shape of the observed positron spectrum from the PAMELA experiment.}    \\begin{document} ",
        "introduction": "The most phenomenologically attractive feature of supersymmetry and in particular the minimal supersymmetric standard model (MSSM) is that it offers a natural candidate for dark matter in terms of the lightest superparticle (LSP) \\cite{1.}. Astrophysical constraints on dark matter requires it to be a neutral and colourless particle, while direct detection experiments strongly disfavour a sneutrino LSP. That makes the lightest neutralino state ${\\tilde{\\chi}}^0_1$ (abbreviated as $\\chi$) the favoured candidate for dark matter in the MSSM.  In the constrained version of the model (CMSSM), corresponding to universal gaugino and scalar masses at the GUT scale, the lightest neutralino state is dominantly a bino over most of the parameter space. Since the bino carries no gauge charge, its main annihilation mechanism is via sfermion exchange in the t-channel. This is usually called the bulk annihilation process; and the region of parameter space giving cosmologically compatible dark matter relic density via this mechanism is called the bulk region. It provides the most natural solution to the dark matter problem, in the sense that the desired dark matter relic density can be obtained in this region with practically no fine-tuning.  However, LEP sets rather stringent lower limits on  the bino LSP as well as the sfermion masses in the CMSSM, which rules out the parameter space of the bulk annihilation region \\cite{2.}.  The reason for the large bino and sfermion mass limits mentioned above is that the LEP lower limit on the neutral Higgs boson mass of the MSSM requires a large radiative correction from top Yukawa coupling, which in turn requires a large stop mass in order to suppress the canceling contribution from stop exchange. This in turn requires a large gluino mass contribution to the RGE of stop mass. Since the GUT scale gluino and bino masses are equal in the CMSSM, this constraint also implies large bino and sfermion masses at the weak scale via their RGE. Evidently a simple way to make the bulk annihilation region of the MSSM dark matter compatible with the Higgs mass limit from LEP is to give up the universality of gaugino masses at the GUT scale; and in particular to assume that the GUT  scale bino mass is significantly smaller than that of gluino. Then the latter can ensure the Higgs mass limit from LEP, while the former ensures relatively small bino and right-handed slepton masses at the weak scale via their RGE, as required for the bulk annihilation region. Moreover there are simple and well motivated models for nonuniversal gaugino masses at the GUT scale, where one assumes that the latter get contributions from SUSY breaking superfields belonging to the nonsinglet representations of the GUT group \\cite{3.,3.5}. One can combine these two observations to construct simple and predictive nonuniversal gaugino mass models, which provide a natural solution to the dark matter relic density while satisfying all the LEP constraints.  The issue of naturalness and fine-tuning involved in achieving the right dark matter relic density \\cite{4.} was investigated in \\cite{5a,5b,6.} for a generic MSSM with nonuniversal gaugino masses. Assuming the usual measure of fine-tuning,  \\begin{equation} \\Delta_a^\\Omega=\\frac{\\partial\\ln (\\Omega_{CDM}h^2)}{\\partial\\ln (a)}\\qquad\\&\\qquad \\Delta^\\Omega=max\\left(\\Delta_a^\\Omega\\right), \\label{relic} \\end{equation} where $a$ refers to the input parameters of the model \\cite{5a}, it was found that $\\Delta^\\Omega\\sim 1$  over the bulk region. This means there is practically no fine-tuning involved in achieving the desired dark matter relic density over the bulk region. In fact over most of this region $\\Delta^\\Omega$ was found to be slightly less than 1, for which the authors called the bulk region 'supernatural' for achieving the desired dark matter relic density. In contrast all the other regions of right dark matter relic density like the stau-coannihilation, the resonant-annihilation and the focus-point regions had 1-2 orders of magnitude higher values of this fine-tuning measure.  Of course one has to pay the usual fine-tuning price for radiative EW symmetry breaking, $\\Delta^{EW}\\sim 100$, for the bulk annihilation region like the other DM relic density compatible regions of the MSSM. However, a quantitative evaluation of this fine-tuning parameter in \\cite{6.} shows, that the bulk annihilation region has one of the lowest $\\Delta^{EW}$ amongst all the DM relic density compatible regions of the MSSM. Thus the low value of $\\Delta^\\Omega$ is achieved here without any additional cost to the $\\Delta^{EW}$.  Subsequently this issue was investigated in a set of simple and predictive nonuniversal gaugino mass models, where the GUT scale gaugino masses are assumed to get contributions from a combination of two SUSY breaking superfields belonging to singlet and nonsinglet representations of the GUT group SU(5) \\cite{7.} - i.e. the combinations 1+24, 1+75 and 1+200.  In each case one could access the bulk region with $\\Delta^\\Omega \\sim 1$, implying practically no fine-tuning required to achieve the right dark matter relic density.   In the present work we have investigated the signatures  of this set of natural SUSY dark matter models for direct and indirect detection experiments. In section 2, we summarize the essential ingredients of the model. In section 3, we present some representative SUSY mass spectra of this model and briefly comment on their implications for the signatures of the model at LHC. Then we present the model predictions for direct and indirect dark matter detection experiments in sections 4 and 5 respectively. In particular we shall see in section 5 that the model predicts a hard positron spectrum like that reported by the PAMELA experiment \\cite{8.}, though it cannot account for the required boost factor in the rate. We conclude with a brief summary of our results in section 6. ",
        "conclusions": "Motivated by the observation that the bulk region of dark matter relic density can be achieved without fine tuning in  models with non-universal gaugino masses at the GUT scale, specifically those arising from a combination of two SUSY breaking superfields belonging to singlet and non-singlet representations of SU(5) \\cite{7.}, we investigate the dark matter phenomenology of these models.  We study the signals of the  1+75 and 1+200 models of \\cite{7.} in direct and indirect detection experiments. We scan the parameter space $M_1^G$=150-250 GeV and $m_0$=80-50 GeV which corresponds to the bulk region and where the bino LSP mass has mass in the range 60-100 GeV. The gluino mass is taken in the range $M_3^G$=600-1000 GeV to evade the bounds on squark and gluino masses from the 7 TeV LHC data \\cite{13.}.  The direct detection cross section for the $\\chi\\, p$ scattering is small for a primarily bino LSP because it is mediated by the higgs which couples to the gaugino and Higgsino components of $\\chi$. The recent Xenon100 result with 225 day exposure rules out DM-proton SI-cross section of up to $2 \\times 10^{-45}$cm$^2$ \\cite{15.}. The 1+75 and 1+200 models studied have a lower $\\chi\\,p$ cross section (Fig 1) and these models are consistent with direct detection experiments so far. A future Xenon 1T experiment which can probe $\\chi\\,p$ cross sections as low as $10^{-47}$ cm$^{2}$ will provide a stringent test of these models.  The dominant process for indirect detection signal of dark matter is via the s-wave radiative annihilation, $\\chi\\chi\\stackrel{\\tilde{l}_R}{\\longrightarrow}\\bar{l}l\\gamma$ \\cite{17.}. This will contribute to the flux of electron/positrons, photons and neutrinos from $\\mu$ and $\\tau$ decays. We find that the 100 GeV bino DM can make a significant contribution to the positron excess observed by PAMELA \\cite{8.}. We see from Fig 2 that the DM annihilation can explain the positron excess (barring the last data point where the signal is lower than the data within 2-sigma) with a boost factor of $\\sim 7000$.  Such a boost factor may be attributed to astrophysical sources\\cite{24.,25.} We do not consider a higher DM mass as that would require obtaining the required relic density by stau-coannihilation which would involve a large fine tuning of the parameters at the GUT scale. Moreover the pair annihilation cross section in the stau-connihilation regime is smaller so a much larger boost factor $\\sim 30,000$ is required \\cite{17.} in order to explain the PAMELA positron signal.  The  measurement of positron flux beyond 100 GeV by AMS2 \\cite{22.} will provide a stringent test of the natural dark matter models \\cite{7.}."
    },
    "1208/1208.4365_arXiv.txt": {
        "abstract": "Intermediate mass black holes (IMBHs) are among the most elusive objects in contemporary astrophysics. Both theoretical and observational evidence of their existence is subject of debate. Conversely, both theory and observations confirm the presence of a large population of millisecond pulsars (MSPs) with low mass companions residing in globular cluster (GC) centers. If IMBHs are common in GC centers as well, then dynamical interactions will inevitably break up many of these binaries, causing the ejection of several fast MSPs in the Galactic halo. Such population of fast halo MSPs, hard to produce with 'standard' MSP generation mechanisms, would provide a strong, albeit indirect, evidence of the presence of a substantial population of IMBHs in GCs. In this paper we study in detail the dynamical formation and evolution of such fast MSPs population, highlighting the relevant observational properties and assessing detection prospects with forthcoming radio surveys. ",
        "introduction": "Black holes (BHs), with masses ranging from tens to billion solar masses, are among the most exciting objects in the astrophysical world. Although there is nowadays plenty of evidence for the existence of stellar mass BHs ($M\\sim 10\\msun$) and massive BHs (MBHs, $M> 10^6\\msun$), there is no convincing observational proof for the existence of BHs in the $100\\msun$-$10^5\\msun$ range, the so called intermediate mass BHs (IMBHs).  IMBHs could form as a result of dynamical instabilities in dense stellar clusters. Despite the large number of works devoted to this problem, any firm conclusion is still lacking. Runaway collisions between massive stars after their sinking into the cluster center provides a channel to form a very massive ($\\sim10^3\\msun$) star in the center. This object would be expected to collapse into an intermediate mass black hole ( Miller \\& Hamilton 2002; Portegies Zwart \\& McMillan 2002; Portegies Zwart et al. 2004) retaining most of the progenitor mass. More detailed simulations including stellar evolution and feedback from stellar winds at solar metallicities, show a systematic inefficiency in growing the central object to masses larger than $\\sim 100\\msun$ (e.g., Glebbeek et al. 2009). IMBHs could then form only in those initially metal poor systems for which stellar mass loss via stellar wind would be strongly reduced.  If theoretical results are unclear, observational evidences are at best elusive.  Globular clusters (GCs), among the densest stellar systems known in galaxies, have become prime sites for IMBH searches. The presence of an IMBH would affect the dynamics of the cluster in several ways, and in the last decade many signatures have been proposed. The presence of an IMBH would 'heat' the stellar density profile, creating a large core (Baumgardt et al. 2005; Trenti 2008; Umbreit et al. 2009). Therefore, IMBHs are likely to reside in those cluster showing a large $r_c/r_h$ ratio, where $r_c$ and $r_h$ are the core and the half light radii, respectively. Observations suggest (Baumgardt et al. 2005) that this may be the case for $\\sim30\\%$ of the Milky Way (MW) GCs. Another clear fingerprint is the Keplerian rising of the velocity dispersion of the stellar distribution inside the sphere of influence of the IMBH. However, for typical IMBH masses, such signature would appear on sub-arcsecond scales, at the limit of current optical facilities. Ibata et al. (2009) report the detection of a Keplerian cusp in M54, consistent with a central $10^4\\msun$ IMBH (although radial anisotropy in the stellar distribution may explain the observations as well). Several other line-of-sight velocity studies were undertaken by Baumgardt et al. (2003a), van den Bosch et al. (2006), and Chakrabarty (2006) on M15, by Gebhardt et al. (2002), Baumgardt et al.  (2003b), Gebhardt et al. (2005) on G1, by Noyola et al. (2008), Sollima et al. (2009), van der Marel \\& Anderson (2010) on Omega Centauri, resulting for now in no undisputed definitive detection. Globular clusters are old systems with relaxation timescales shorter than their lifetimes. The most massive stars should then have had enough time to segregate into the center via dynamical friction. This mass segregation should be partially suppressed by the presence of an IMBH. Several studies on the radial dependence of the stellar mass function in different cluster might be consistent with the presence of an IMBH (see Beccari et al. 2010 for M10, Pasquato et al. 2009 for NGC 2298, Umbreit et al.  2009 for NGC5694). However, the presence of primordial binaries can mimic the same effect of an IMBH, and it is not clear to what extent the results depend on the chosen initial conditions. Lastly, an IMBH would accrete ambient gas, and such accretion activity may be observable in X-ray and/or radio.  Measurements of radio luminosity of few Galactic GCs put stringent upper limits on IMBH masses (Maccarone \\& Servillat 2008). Similarly, Cseh et al. (2010) impose a limit of 1500$\\msun$ to any putative massive object in the center of NGC 6388, by using combined X-ray and radio observations. We notice, however, that these limits are affected by a number of parameters (like the accretion efficiency, the gas content in GCs, etc.) that are not constrained themselves.  Millisecond pulsars (MSPs) are usually associated with dense stellar environment, as their formation channel requires a recycling phase within a binary (Camilo \\& Rasio 2006, Lorimer 2008): almost two thirds (140 out of 220) of the known MSPs are observed in GCs, most of them in binary systems residing in the GC cores (Lorimer 2008). The simultaneous presence of an IMBH and of MSPs in binary systems in the same GC opens interesting dynamical possibilities.  Dynamical interactions with an IMBH are expected to break-up binaries ejecting one of the components while leaving the other one bound to the IMBH (Hills 1988, see Devecchi et al. 2007 in the context of IMBHs). One consequence of the presence of IMBHs in GCs would then be the presence of a halo population of fast MSPs, ejected by this tidal break-up process. The detection of such population of fast halo MSPs would not provide a direct evidence for the presence of an IMBH in any particular cluster, but will prove the existence of a population of IMBHs residing in GCs and capable of ejecting MSPs via dynamical interactions. In this paper, we investigate the properties of the expected fast halo MSP population for a range of formation models, assessing observability with forthcoming radio instruments, such as the Square Kilometre Array (SKA, Lazio 2009).  The paper is organized as follows. In Section 2 we introduce the main ingredients of our models, namely the GC IMBH and MSP populations. The binary-IMBH dynamical interaction is detailed in Section 3. In Section 4 we present the population of ejected MSPs highlighting their relevant properties and we discuss observational prospects in Section 5. Our main findings are summarized in Section 6. ",
        "conclusions": "Inspired by the detection of several MSPs in binaries in the core of Galactic GCs, we investigated their dynamical interaction with putative IMBHs lurking in the center of the same GCs. As a consequence of three body interactions, MSPs are ejected in the Galactic halo with velocities (with respect to the host GC) up to several hundreds km s$^{-1}$, sufficient to distribute them into the outer halo, but not to escape the Galaxy potential, allowing the formation of a substantial population of fast halo MSPs. The detection of such population would be a strong element, albeit indirect, in favor of the presence of IMBHs in the center of GCs.  We investigated four ejection scenarios, involving two different IMBH populations and MSPs formation histories. We considered a high -H- and a low -L- IMBH mass functions, consistent with the runaway model of Portegies Zwart et al. (2005) and with the low mass end extrapolation of the $M-\\sigma$ relation respectively. We also assumed two different MSPs formation scenarios, one in which MSPs are simultaneously generated in a single burst 12Gyr ago, and one with a constant MSP formation rate along the whole Galaxy history. In both cases, the MSP population in GCs was normalized to match the observed numbers of MSPs in close binaries in the clusters Terzan 5 and 47 Tuc, and then scaled to all other GCs according to their structural parameters. The investigated scenarios predict between 100 and 600 fast MSPs wandering in the MW halo, a population that can in principle be detected with future radio surveys. We emphasize here that all the considered models are consistent with current MSP observations in GCs; that is, even considering efficient binary break-up and MSP ejection, GCs retain a substantial population of binaries containing MSPs, as observed today.  In the spirit of using fast halo MSPs as a probe of IMBHs in GCs, we checked that such population cannot result from alternative channels. Halo MSPs can be in principle produced by other dynamical ejection mechanisms, such as binary-binary or three body interaction with stellar BHs, or can be relics of the natural evolution of the old halo stellar population. Both channels looks inefficient in producing a sizable population of halo MSPs, unless several tens of stellar BHs are retained in GCs, or the MSP formation efficiency is much higher in the halo than in the disk.  We finally ran fully consistent simulations of the MSP population in the Galaxy. We considered 30000 MSPs distributed in the disk plus several fast halo MSP populations predicted by our models. The disk and the halo populations were averaged over 1000  and 20 realizations respectively. We simulated observations with SKA, taking into account for selection effects and limiting sensitivity of the instrument. We found the distribution of projected distances from the galactic plane $h$ to be an effective discriminant between the two populations. Fast halo MSPs, in fact, show a significant excess of objects at high $h$, in sharp contrast with standard disk-born MSPs. This means that any excess of MSPs at high $h$  would be a strong hint of the presence of a different MSP population. Here we showed that MSP ejection from GC, due to dynamical interactions with IMBHs, would easily produce such excess, and can therefore leave a clear imprint in future radio surveys. We quantified the statistical significance of halo MSP detection by performing tests on synthetic samples drawn from the disk plus halo distribution. Results shown in figure \\ref{f8} and table \\ref{tab2} are very encouraging. If we isolate the distribution of MSPs observed at distances from the galactic plane larger than $h_{\\rm cut}=3$ kpc, the halo population is, on average, detected at a 97\\% confidence level for all models but L-I-${\\cal A}$ and L-I-${\\cal B}$. When $h_{\\rm cut}=4$ kpc, detection confidence is boosted to better than 99\\% for all models. This is also true for the reduced SKA ${\\cal B}$ configuration, indicating that such detection will be easily feasible.   Our analysis is admittedly oversimplified in several respects. Most importantly, the disk MSP population is hard to model, since $\\sim100$ objects has been detected to date. This introduces considerable uncertainties in our results. Note however that, unless disk MSPs are far more numerous than suggested by present observation or extend far above the MW disk height, our basic result, that an excess of MSPs located at large distances from the Galactic plane would strongly support dynamical ejection due to IMBH residing in GCs, remains unaltered."
    },
    "1208/1208.6010_arXiv.txt": {
        "abstract": "Strong gravitational lenses with measured time delays between the multiple images and models of the lens mass distribution allow a one-step determination of the time-delay distance, and thus a measure of cosmological parameters.  We present a blind analysis of the gravitational lens \\rxj\\ incorporating (1) the newly measured time delays from COSMOGRAIL, the COSmological MOnitoring of GRAvItational Lenses, (2) archival \\textit{Hubble Space Telescope} imaging of the lens system, (3) a new velocity-dispersion measurement of the lens galaxy of $323\\pm20\\,\\kms$ based on Keck spectroscopy, and (4) a characterization of the line-of-sight structures via observations of the lens' environment and ray tracing through the Millennium Simulation.  Our blind analysis is designed to prevent experimenter bias. The joint analysis of the data sets allows a time-delay distance measurement to 6\\% precision that takes into account all known systematic uncertainties.  In combination with the \\textit{Wilkinson Microwave Anisotropy Probe} seven-year (WMAP7) data set in flat $w$CDM cosmology, our unblinded cosmological constraints for \\rxj\\ are: $H_0=80.0^{+5.8}_{-5.7}\\,\\kmsMpc$, $\\Ode=0.79\\pm0.03$, $w=-1.25^{+0.17}_{-0.21}$.  We find the results to be statistically consistent with those from the analysis of the gravitational lens \\blens, permitting us to combine the inferences from these two lenses. The joint constraints from the two lenses and WMAP7 are $H_0=75.2^{+4.4}_{-4.2}\\,\\kmsMpc$, $\\Ode=0.76^{+0.02}_{-0.03}$ and $w = -1.14^{+0.17}_{-0.20}$ in flat $w$CDM, and $H_0=73.1^{+2.4}_{-3.6}\\,\\kmsMpc$, $\\OL=0.75^{+0.01}_{-0.02}$ and $\\Ok=0.003^{+0.005}_{-0.006}$ in open $\\Lambda$CDM.  Time-delay lenses constrain especially tightly the Hubble constant $H_0$ (5.7\\% and 4.0\\% respectively in $w$CDM and open $\\Lambda$CDM) and curvature of the universe. The overall information content is similar to that of Baryon Acoustic Oscillation experiments. Thus, they complement well other cosmological probes, and provide an independent check of unknown systematics.  Our measurement of the Hubble constant is completely independent of those based on the local distance ladder method, providing an important consistency check of the standard cosmological model and of general relativity. ",
        "introduction": "\\label{sec:intro}  In the past century precise astrophysical measurements of the geometry and content of the universe (hereafter cosmography) have led to some of the most remarkable discoveries in all of physics. These include the expansion and acceleration of the Universe, its large scale structure, and the existence of non-baryonic dark matter \\citep[see review by][]{FreedmanTurner03}. These observations form the empirical foundations of the standard cosmological model, which is based on general relativity and the standard model of particle physics but requires additional non-standard features such as non-baryonic dark matter and dark energy.  Even in the present era of so-called precision cosmography, many profound questions about the Universe remain unanswered. What is the nature of dark energy? What are the properties of the dark matter particle? How many families of relativistic particles are there? What are the masses of the neutrinos? Is general relativity the correct theory of gravity? Did the Universe undergo an inflationary phase in its early stages?  From an empirical point of view, the way to address these questions is to increase the accuracy and precision of cosmographic experiments. For example, clues about the nature of dark energy can be gathered by measuring the expansion history of the Universe to very high precision, and modeling the expansion as being due to a dark energy component having an equation of state parametrized by $w$ that evolves with cosmic time \\citep[\\eg][and references therein]{FTH08}. Likewise, competing inflationary models can be tested by measuring the curvature of the Universe to very high precision.  Given the high stakes involved, it is essential to develop multiple independent methods as a way to control for known systematic uncertainties, uncover new ones, and ultimately discover discrepancies that may reveal new fundamental physics. For example, a proven inconsistency between inferences at high redshift from the study of the cosmic microwave background, with inferences at lower redshift from galaxy redshift surveys would challenge the standard description of the evolution of the Universe over this redshift interval, and possibly lead to revisions of either our theory of gravity or of our assumptions about the nature of dark matter and dark energy.  In this paper we present new results from an observational program aimed at precision cosmography using gravitational lens time delays. The idea of doing cosmography with time-delay lenses goes back fifty years and it is a simple one \\citep{Refsdal64}. When a source is observed through a strong gravitational lens, multiple images form at the extrema of the time-delay surface, according to Fermat's principle \\citep[e.g.,][]{SEF92, Falco05,SchneiderEtal06}. If the source is variable, the time delays between the images can be measured by careful monitoring of the image light curves \\citep[see, \\eg][]{Courbin03}. With an accurate model of the gravitational lens, the absolute time delays can be used to convert angles on the sky into an absolute distance, the so-called time-delay distance, which can be compared with predictions from the cosmological model given the lens and source redshifts \\citep[\\eg][and references therein]{BlandfordNarayan92, Jackson07, Treu10}. This distance is a combination of three angular diameter distances, and so is primarily sensitive to the Hubble constant ($H_0$), with some higher order dependence on the other cosmological parameters \\citep{CoeMoustakas09,Linder11}. Gravitational time delays are a one-step cosmological method to determine the Hubble constant that is completely independent of the local cosmic distance ladder \\citep{FreedmanEtal01,RiessEtal11, FreedmanEtal12, ReidEtal12}. Knowledge of the Hubble constant is currently the key limiting factor in measuring parameters like the dark energy equation of state, curvature, or neutrino mass, in combination with other probes like the cosmic microwave background \\citep{FreedmanMadore10,RiessEtal11, FreedmanEtal12, WeinbergEtal12, SuyuEtal12b}.  These features make strong gravitational time delays a very attractive probe of cosmology.  Like most high-precision measurements, however, a good idea is only the starting point. A substantial amount of effort and observational resources needs to be invested to control the systematic errors. In the case of gravitational time delays, this has required several observational and modeling breakthroughs. Accurate, long duration, and well-sampled light curves are necessary to obtain accurate time delays in the presence of microlensing. Modern light curves have much higher photometric precision, sampling and duration \\citep{FassnachtEtal02,CourbinEtal11} compared to the early pioneering light curves \\citep[\\eg][]{LeharEtal92}. High resolution images of extended features in the source, and stellar kinematics of the main deflector, provide hundreds to thousands of data points to constrain the mass model of the main deflector, thus reducing the degeneracy between the distance and the gravitational potential of the lens that affected previous models constrained only by the positions of the lensed quasars \\citep[e.g.,][]{SchechterEtal97}. Finally, cosmological numerical simulations can now be used to characterize the distribution of mass along the line of sight (LOS) \\citep{HilbertEtal09}, which was usually neglected in early studies that were not aiming for precisions of a few percent. The advances in the use of gravitational time delays as a cosmographic probe are summarized in the analysis of the gravitational lens system \\blens\\ by \\citet{SuyuEtal10}. In that paper, we demonstrated that, with sufficient ancillary data, a single gravitational lens can yield a time-delay distance measured to 5\\% precision, and the Hubble constant to 7\\% precision. In combination with the \\textit{Wilkinson Microwave Anisotropy Probe} 5-year (WMAP5) results, the \\blens\\ time-delay distance constrained $w$ to 18\\% precision and the curvature parameter to $\\pm0.02$ precision, comparable to contemporary Baryon Acoustic Oscillation experiments \\citep{PercivalEtal07} and observations of the growth of massive galaxy clusters \\citep[on $w$ constraints;][]{MantzEtal10}.  Building on these recent developments in the analysis, and on the state-of-the-art monitoring campaigns carried out by the COSMOGRAIL (COSmological MOnitoring of GRAvItational Lenses; e.g., \\citeauthor{VuissozEtal08}~\\citeyear{VuissozEtal08}; \\citeauthor{CourbinEtal11}~\\citeyear{CourbinEtal11}; \\citeauthor{TewesEtal12b}~\\citeyear{TewesEtal12b}) and \\citet{KochanekEtal06} teams, it is now possible to take gravitational time delay lens cosmography to the next level and achieve precision comparable to current measurements of the Hubble constant, flatness, $w$ and other cosmological parameters \\citep{RiessEtal11, FreedmanEtal12, KomatsuEtal11}. To this end we have recently initiated a program to obtain data and model four additional gravitational lens systems with the same quality as that of \\blens.  These four lenses are selected from the COSMOGRAIL sample with the smallest uncertainties in the delays between the images of $\\leq$$6\\%$.  They cover various image configurations: (1) four lensed images with three of them merging, a.k.a.~the ``cusp'' configuration, (2) four lensed images with two of them merging, a.k.a.~the ``fold'' configuration, (3) four images that are nearly symmetric about the lens center, a.k.a.~the ``cross'' configuration, and (4) two images on opposite sides of the lens galaxy.  This sample will allow us to probe the optimal lens configuration for time-delay cosmography and also investigate potential selection effects.  We present here the results for the first of these systems, \\rxj, based on new time delays measured by the COSMOGRAIL collaboration \\citep{TewesEtal12b}, new spectroscopic data from the Keck Telescope, a new analysis of archival \\textit{Hubble Space Telescope} (\\hst) images, and a characterization of the LOS effects through numerical simulations, the observed galaxy number counts in the field, and the modeled external shear. We carry out a self-consistent modeling of all the available data sets in a Bayesian framework, and infer (1) a likelihood function for the time-delay distance that can be combined with any other independent probe of cosmology, and (2) in combination with our previous measurement of \\blens\\ and the WMAP 7-year (WMAP7) results, the posterior probability density  function (PDF) for the Hubble constant, curvature density parameter and dark energy equation-of-state parameter $w$.  Three additional lens systems are scheduled to be observed with \\hst\\ in cycle 20 (GO 12889; PI Suyu) and will be published in forthcoming papers. An integral part of this program is the use of blind analysis, to uncover unknown systematic errors and to avoid unconscious experimenter bias. Only when each system's analysis has been judged to be complete and final by its authors, are the implications for cosmology revealed. These results are then published without any further modification. In this way, we can assess whether the results are mutually consistent within the estimated errors or whether unknown systematics are adding significantly to the total error budget.  This paper is organized as follows. After a brief recap of the theory behind time-delay lens cosmography in \\sref{sec:theory}, we summarize our strategy in \\sref{sec:strgy} and describe our observational data in \\sref{sec:obs}. In \\sref{sec:prob}, we write out the probability theory used in the data modeling and describe the procedure for carrying out the blind analysis.  The lensing and time-delay analysis are presented in \\sref{sec:lensmod}, and a description of our treatment of the LOS mass structure in the \\rxj\\ field is in \\sref{sec:breakmsd}.   We present measurements of the time-delay distance, and discuss the sources of uncertainties in \\sref{sec:tdist}.  We show our unblinded cosmological parameter inferences in \\sref{sec:cosmo}, which includes joint analysis with our previous lens data set and with WMAP7. Finally, we conclude in \\sref{sec:conclude}.  Throughout this paper, each quoted parameter estimate is the median of the appropriate one-dimensional marginalized posterior PDF, with the quoted uncertainties showing, unless otherwise stated, the $16^{\\rm th}$ and $84^{\\rm th}$ percentiles (that is, the bounds of a 68\\% credible interval). ",
        "conclusions": " \\begin{enumerate} \\item Our comprehensive lens model reproduces the global features of the \\hst\\ image and the time delays.  We quantify the uncertainty due to the deflector gravitational potential on the time-delay distance to be at the $4.6\\%$ level. \\item Based on the external shear strength from the lens model and the overdensity of galaxy count around the lens, we obtained a PDF for the external convergence by ray tracing through the Millennium Simulation.  This $\\kext$ PDF contributes to the uncertainty on $\\tdist$ also at the $4.6\\%$ level. \\item Our robust time-delay distance measurement of $6\\%$ takes into account all sources of known statistical and systematic uncertainty.  We provide a fitting formula to describe the PDF of the time-delay distance that can be used to combine with any other independent cosmological probe. \\item The time-delay distance of \\rxj\\ is mostly sensitive to $H_0$, especially given the low redshift of the lens. \\item Assuming a flat $\\Lambda$CDM with fixed $\\OL=0.73$ and uniform prior on $H_0$, our unblinded $H_0$ measurement from \\rxj\\ is $78.7^{+4.3}_{-4.5}\\,\\kmsMpc$. \\item The constraint on $H_0$ helps break parameter degeneracies in the CMB data. In combination with WMAP7 in $w$CDM, we find $H_0=80.0^{+5.8}_{-5.7}\\,\\kmsMpc$, $\\Ode=0.79\\pm0.03$, and $w=-1.25^{+0.17}_{-0.21}$.  These are statistically consistent with the results from the gravitational lens \\blens.  There are no significant residual systematics detected in our method based on this combined analysis of the two systems. \\item By combining \\rxj, \\blens\\ and WMAP7, we derive the following constraints: $H_0=75.2^{+4.4}_{-4.2}\\,\\kmsMpc$, $\\Ode=0.76^{+0.02}_{-0.03}$ and $w = -1.14^{+0.17}_{-0.20}$ in flat $w$CDM, and $H_0=73.1^{+2.4}_{-3.6}\\,\\kmsMpc$, $\\OL=0.75^{+0.01}_{-0.02}$ and $\\Ok=0.003^{+0.005}_{-0.006}$ in open $\\Lambda$CDM. \\item Our measurement of the Hubble constant is completely independent of those based on the local distance ladder method \\citep[e.g.,][]{RiessEtal11, FreedmanEtal12}, providing an important consistency check of the standard cosmological model and of general relativity. \\item A comparison of the lenses and other cosmological probes that are each combined with WMAP7 shows that the constraints from the lenses are comparable in precision to various state-of-the-art probes.  Lenses are particularly powerful in measuring the spatial curvature of the universe, and are complementary to other cosmological probes. \\end{enumerate}  Thanks to the dedicated monitoring by the COSMOGRAIL \\citep[e.g.,][]{VuissozEtal08, CourbinEtal11, TewesEtal12b, TewesEtal12a} and \\citet{KochanekEtal06} collaborations, the number of lenses with accurate and precise time delays are increasing. Deep \\hst\\ imaging for three of these lenses will be obtained in cycle 20 to allow accurate lens mass modeling that turns the delays into distances. Using the estimated uncertainties of the time-delay distances of the three lenses, we expect to measure $H_0$ from our assembled sample of five lenses (\\blens, \\rxj, and the three cycle 20 lenses) to roughly $3.8\\%$ in a $w$CDM cosmology if no significant residual systematics are detected. Current and upcoming telescopes and surveys including the Panoramic Survey Telescope \\& Rapid Response System, Hyper-Suprime Camera on the Subaru Telescope, and Dark Energy Survey expect to detect hundreds of AGN lenses with dozens of delays measured \\citep{OguriMarshall10}.  Ultimately, the Large Synoptic Survey Telescope will discover thousands of time-delay lenses, painting a bright future for cosmography with gravitational lens time delays."
    },
    "1208/1208.1406_arXiv.txt": {
        "abstract": "The origin and properties of the cosmic radiation are one of the most intriguing question in modern astrophysics. The precise measurement of the chemical composition and energy spectra of the cosmic rays provides fundamental insight into these subjects. In this paper we will review the existing experimental data. Specifically, we will analyse results collected by space-born experiments discussing the experimental uncertainties and challenges with a focus on the PAMELA experiment. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.6226_arXiv.txt": {
        "abstract": "We have analyzed a sample of hot subdwarfs (sdB, sdO) selected from the {\\it GALEX} ultraviolet sky survey. Applying a model atmosphere analysis we determined the temperature, surface gravity, and helium-to-hydrogen abundance ratio, and obtained preliminary constraints on the CNO abundance for a sample of 181 stars. Adopting colourimetric (ultraviolet-infrared) and quantitative spectral decomposition we also investigated the incidence of solar type or earlier (A, F, G) companions. ",
        "introduction": "Hot subdwarfs are core He burning stars located at the blue end of the horizontal branch (HB) or extreme horizontal branch (EHB). These subluminous objects are just below the early-type main-sequence (MS) stars in the Hertzsprung-Russell diagram (HRD). Being in a relatively long lasting ($\\sim$160 Myr) intermediate evolutionary stage of $\\sim$1$M_{\\odot}$ stars they are quite common and are the primary sources of the UV excess of elliptical galaxies, and overwhelm white dwarfs in blue and UV surveys of old stellar populations.   To look for bright, thus nearby white dwarf candidates in the {\\it GALEX} database, \\citet{vennes11a} devised a method based on UV, optical and infrared colour criteria and sorted out $\\sim$200 bright ($N_{\\rm {\\small UV}}$$<$14) and hot stars. Between 2008 and 2011, systematic, low-resolution spectroscopic follow-up with the ESO/NTT and NOAO/Mayall telescopes confirmed 167 stars as hot subdwarfs. Here, we present the results of a model atmosphere analysis of 127 sdB and 40 sdO stars. Such large and homogeneously modelled samples of bright subdwarfs are useful in identifying candidates for pulsation (e.g., \\citealp{ostensen10}) and radial velocity studies (e.g., \\citealp{geier11}). ",
        "conclusions": "We have presented a subdwarf sample of 127 sdB and 40 sdO stars from the GALEX survey. Our selection provides a complete, homogeneously modelled sample of hot subdwarf stars, similar to the HS and SPY surveys. This high-quality set confirms well-known correlations between effective temperature and He abundance. A dichotomy seems to emerge for sdB stars that separates stars around 28000 K and $\\log g=5.45$ from stars near 33500 K and $\\log g=5.8$. This separation is seen in the He abundance pattern and its connection to binarity will be investigated further. The full catalogue, along with further details of spectral modelling, binary decomposition and on several individual stars will be presented in a forthcoming paper."
    },
    "1208/1208.1267.txt": {
        "abstract": "Radiative transfer studies of Type Ia supernovae (SNe\\,Ia) hold the promise of constraining both the density profile of the SN ejecta and its stratification by element abundance which, in turn, may discriminate between different explosion mechanisms and progenitor classes. Here we analyse the Type Ia SN\\,2010jn (PTF10ygu) in detail, presenting and evaluating near-ultraviolet (near-UV) spectra from the \\textit{Hubble Space Telescope} and ground-based optical spectra and light curves. SN\\,2010jn was discovered by the Palomar Transient Factory (PTF) 15 days before maximum light, allowing us to secure a time series of four near-UV spectra at epochs from $-$10.5 to $+$4.8 days relative to $B$-band maximum. The photospheric near-UV spectra are excellent diagnostics of the iron-group abundances in the outer layers of the ejecta, particularly those at very early times.  Using the method of `Abundance Tomography' we derive iron-group abundances in SN\\,2010jn with a precision better than in any previously studied SN\\,Ia. Optimum fits to the data can be obtained if burned material is present even at high velocities, including significant mass fractions of iron-group elements.  This is consistent with the slow decline rate (or high `stretch') of the light curve of SN\\,2010jn, and consistent with the results of delayed-detonation models. Early-phase UV spectra and detailed time-dependent series of further SNe\\,Ia offer a promising probe of the nature of the SN\\,Ia mechanism. ",
        "introduction": "Supernovae (SNe) play an important role in many areas of modern astrophysics. In particular, Type Ia SNe (SNe\\,Ia) produce most of the iron-group elements in the cosmos \\citep{iwa99}, and can be used as `standardizable candles' to probe the expansion history of the Universe \\citep[\\eg][]{rie98,per99,rie07,kes09,sul11,suz12}. Because of their importance, extensive efforts are under way to build a comprehensive observational and theoretical picture of SNe\\,Ia (\\eg \\citealt{hil00,maz07,howell11,roepke11}). A primary aim is to understand the nature and diversity of the explosions and thus to place on a physical ground the empirical calibration procedures that are used to deduce the luminosity of SNe\\,Ia from their light-curve shape \\citep[\\eg][]{phi93,phi99}. Another aim is to determine the progenitor systems of SNe\\,Ia.  One way to improve our understanding of SN\\,Ia physics is to analyse and interpret detailed observations through radiation transport models. A particularly useful technique is that of `Abundance Tomography' \\citep{ste05}, where a time series of SN\\,Ia spectra is modelled and information obtained about the density profile of the SN ejecta and its stratification in terms of element abundances. This makes it possible to describe the mode of explosion and possibly to discriminate among different progenitor scenarios. Many detailed spectral time series for SNe\\,Ia have become available in the last decade, but most of them are restricted to optical wavelengths.  The ultraviolet (UV) spectrum of SNe\\,Ia is shaped by iron-group elements, which dominate line-blocking and fluorescence effects \\citep{kirshner93,pauldrach96,maz00}. Thus, series of UV spectra are invaluable diagnostics for iron-group abundances and ejecta densities in different zones of the SN.  These in turn are indicators for basic explosion properties such as burning efficiency or explosion energy \\citep{iwa99}. Extensive work on UV spectra of SNe\\,Ia has been carried out especially in the last few years. On the observational side, compilations of UV spectra have been built up and studied \\citep[\\eg][]{ellis08a,fol08,buf09,fol12sdss,maguire12a}, but also the properties of individual objects have been investigated in more detail \\citep[\\eg][]{fol12sn2009ig}. On the theoretical side, spectral modelling has been performed to understand the diagnostic utility of UV spectra better \\citep[\\eg][]{sau08,fol12sn2011iv,wal12}. Both observations and models indicate that even SNe\\,Ia with similar optical spectra can be very different in the UV \\citep{lentz00a, ellis08a, sau08, wal12}.  Here we present observations and models of the SN 2010jn (PTF10ygu), a `normal' SN\\,Ia \\citep[\\cf][]{bra93}, which is however quite luminous and has high line velocities as well as high-velocity features (HVFs, \\eg \\citealt{mazzali05a}) in the spectra. The SN was discovered by the Palomar Transient Factory (PTF) only a few days after explosion, which allowed a detailed spectral time series to be obtained, including near-UV data from the \\textit{Hubble Space Telescope} (\\textit{HST}). The series of combined optical-UV spectra made it possible to analyse the outer and intermediate ejecta in unprecedented detail with the tomography technique. The abundance stratification was inferred based on density profiles of single-degenerate Chandrasekhar-mass explosion models. These models \\citep[\\cf \\eg][]{hil00} assume that SNe\\,Ia are explosions of accreting carbon-oxygen (CO) white dwarfs (WDs). We used both the fast-deflagration model W7 \\citep{nom84w7,iwa99}, which has been shown to match average SNe\\,Ia \\citep{ste05,tanaka11}, and a more energetic delayed-detonation model (WS15DD3, \\citealt{iwa99}). The latter model assumes that the combustion flame, which initially propagates subsonically (deflagration), becomes supersonic (detonation) at some point \\citep{kho91a}. Based on the quality of the fits, we suggest that the latter model is a more realistic description of SN\\,2010jn.  The paper starts with a report on the observations and a presentation of the observed spectra, which were obtained from \\mbox{$-\\textrm{13}$\\,{}d} to $+\\textrm{5}$\\,{}d (time in the SN rest frame) with respect to maximum light in the rest-frame $B$ band (Section \\ref{sec:observations}). Afterwards, we discuss our models. We lay out the objectives, methods and assumptions (Sections \\ref{sec:models-objectives} -- \\ref{sec:models-assumptions}). Synthetic spectra are presented and the inferred abundance profiles are discussed (Sections \\ref{sec:tomography-10jn-w7}, \\ref{sec:tomography-10jn-wdd3}). Finally, the results are summarised and conclusions are drawn (Section \\ref{sec:conclusions}). ",
        "conclusions": "\\label{sec:conclusions}  We have for the first time used a time-series of photospheric near-UV\\,/\\,optical spectra of a SN\\,Ia (SN\\,2010jn\\,/\\,PTF10ygu) to analyse its abundance stratification with the tomography technique of \\citet{ste05}. Our work can be seen in line with classical work aimed at understanding UV spectra of SNe Ia \\citep[\\eg][]{kirshner93,maz00}, with significantly refined methods for radiative transfer and spectral modelling. The early near-UV observations have proved extremely valuable for analysing the iron-group content in the outer layers. They allowed us to differentiate among different iron-group elements present in the partially burned and unburned zone. The later spectra give constraints on the abundances further inwards.  From our abundance analysis, we find that SN\\,2010jn synthesised significant amounts of iron-group elements, typical for an energetic, luminous SN Ia. SN\\,2010jn has only a thin oxygen-dominated zone and a limited IME zone. Even these zones contain directly synthesised Fe ($\\sim \\textrm{0.01}$\\,\\Msun) as well as \\Nifs\\ and decay products ($\\sim \\textrm{0.05}$\\,\\Msun\\ in total). The presence of iron-group elements and intermediate-mass elements up to high velocities (even $>$\\,20000\\,\\kms) enhances the opacity of the SN, which is consistent with the overall slow light-curve evolution.  We favour for SN\\,2010jn a Chandrasekhar-mass delayed-detonation model with efficient nucleosynthesis and an explosion energy somewhat above average (WDD3, \\citealt{iwa99}). It allows us to reproduce the high expansion velocities in the observed spectra better than the `fast deflagration' model W7 \\citep{nom84w7,iwa99}, which has little material at high velocities. The presence of Fe and \\Nifs\\ in the outermost layers is not predicted by the original (1D) WDD3 model. The abundance of Fe at high velocities (\\mbox{$\\textrm{13350}\\textrm{\\,\\kms} \\leq v \\leq \\textrm{33000}$\\,\\kms}) is between one and two orders of magnitude above solar. This cannot come from the progenitor; it is rather a clue about the explosion properties of SN\\,2010jn. A high iron-group abundance is consistent with an energetic SN, but still, explaining iron-group material in the outer layers is a challenge for explosion models. Multi-dimensional models may show some clumpiness or asymmetry which could appear to an observer as `outwards-mixing' of iron-group material.  Our models favour a longer rest-frame rise time for SN\\,2010jn ($\\sim$$\\textrm{20}$\\,d) than predicted by the empirical $t^2$ analysis of the early light curve (18.6\\,d). This may imply that the SN after explosion has a `dark phase' longer than predicted by the $t^2$ (`fireball') model. It remains to be seen whether this interesting `quirk' is found in other objects that are not as luminous as SN\\,2010jn. Very early data are best suited to tackle this issue.  Finally, we remark that a precise analysis of the abundances in the outer layers, as presented here, should be complemented by an analysis of the ejecta core, for which nebular spectra are necessary. Also, observations of more SNe\\,Ia in the UV will be needed to establish the properties of the outermost layers of SNe\\,Ia with different characteristics and light-curve properties."
    },
    "1208/1208.0827_arXiv.txt": {
        "abstract": "Dark matter particles captured by the Sun through scattering may annihilate and produce neutrinos, which escape.  Current searches are for the few high-energy neutrinos produced in the prompt decays of some final states.  We show that interactions in the solar medium lead to a large number of pions for nearly all final states.  Positive pions and muons decay at rest, producing low-energy neutrinos with known spectra, including $\\bar{\\nu}_e$ through neutrino mixing.  We demonstrate that Super-Kamiokande can thereby provide a new probe of the spin-dependent WIMP-proton cross section.  Compared to other methods, the sensitivity is competitive and the uncertainties are complementary. ",
        "introduction": "If dark matter is a thermal relic of the early universe, then its self-annihilation cross section is revealed by its present mass density.  To match observations, the required cross section, averaged over relative velocities, is $\\langle \\sigma_A v \\rangle = (5.2 - 2.2) \\times 10^{-26}$~cm$^{3}$/s, as a function of increasing mass, $m_\\chi$~\\cite{Steigman:2012nb}.  This indicates a weakly interacting massive particle (WIMP; denoted $\\chi$)~\\cite{Jungman:1995df, Bertone:2004pz, Feng:2010gw}.  The total annihilation cross section, including all final states, is well defined, but the partial annihilation, scattering, and production cross sections with any specific standard model (SM) particles are model-dependent.  Measurement of any of these cross sections would dramatically constrain WIMP models and eliminate more exotic possibilities.  There are limits from straightforward searches for astrophysical fluxes of annihilation products, direct nuclear scatterings in underground experiments, and collider events with missing energy.  Each of these searches has different underlying uncertainties and one technique may be much more sensitive than others if the WIMP framework is different than commonly supposed.  A convincing WIMP discovery will require observations by multiple experiments with different techniques that give a consistent picture. To gain detailed knowledge of WIMP properties and distributions, observations in multiple channels with comparable sensitivities are needed.  Searches for high-energy neutrinos from the Sun give strong limits on WIMP-nucleon scattering.  WIMPs passing through the Sun may rarely scatter with nuclei and become gravitationally bound.  Scattering can occur by coherent spin-independent~(SI) or valence spin-dependent~(SD) interactions.  Further scatterings thermalize WIMPs in the solar core, where they annihilate.  Only neutrinos can escape and potentially be detected.  When the capture rate, $\\Gamma_{\\rm C}$, and the annihilation rate, $\\Gamma_{\\rm A}$, are in equilibrium, as expected, an upper limit on the neutrino flux sets an upper limit on the WIMP-nucleon scattering cross section.  The most interesting limits using searches for high-energy neutrinos from the Sun are on SD WIMP-proton scattering, $\\sigma_{\\chi p}^{\\rm SD}$.  Though the searches are based on the annihilation process, these limits are independent of $\\langle \\sigma_A v \\rangle$, except for assumptions about the annihilation final states, which govern the detectability of the neutrinos.  High-energy neutrinos come from the few annihilation products that decay promptly, before losing energy in the solar medium, giving continuum spectra up to $E_\\nu \\sim m_\\chi$ (direct annihilation to neutrino pairs, helicity-suppressed in many models, produces a line at $E_\\nu = m_\\chi$).  Strong limits on the SD WIMP-proton scattering cross section have been derived for large $m_\\chi$ and certain final states, such as $W^+W^-, \\tau^+\\tau^-,$ and $b \\bar{b}$.  Generalizing to other final states and a broader range of masses is challenging.  The number of high-energy neutrinos per annihilation is small ($N_\\nu \\sim 1$ for favorable final states) and their spectrum depends on the unknown final state.  Neutrino signal detection and atmospheric neutrino background rejection become easier for larger $m_\\chi$, though neutrinos with $E_\\nu \\gtrsim 100$~GeV are significantly attenuated by interactions in the Sun.  What about the more common but seemingly less favorable final states?  As is well known, most final states ultimately produce pions and muons, which quickly lose energy and decay at rest, producing only MeV neutrinos, which have long been considered undetectable (e.g., see Ref.~\\cite{Gaisser:1986ha, Ritz:1987mh} and many subsequent papers).  We propose a new probe of SD WIMP-proton scattering in standard solar WIMP capture scenarios.  First, we show that the pion yield from WIMP annihilation in the Sun -- produced directly in hadronic decays and further through inelastic interactions in the dense medium -- is large, relatively model-independent, and increases as $N_\\pi \\propto m_\\chi$.   Second, we show that the subsequent low-energy neutrinos are much more detectable than previously thought, due to their high yield ($N_\\nu \\gg 1$) and known spectra, the low atmospheric neutrino backgrounds at low energy, and advances in detectors.  The advantages of this method are sensitivity to low WIMP masses and near-insensitivity to the choice of final state.  While the signal in high-energy neutrinos could vanish for unfavorable annihilation channels, for example if WIMPs annihilated only to light quarks, the low-energy signals would not.  In the following, we review solar WIMP capture and annihilation, calculate the pion and neutrino yields per annihilation, estimate the signal and background rates in Super-Kamiokande (Super-K)~\\cite{Fukuda:2002uc}, and derive new constraints on the SD WIMP-proton cross section.  We conclude by discussing likely improvements and the importance of complementary methods to allow for possible surprises in the astrophysics or physics of dark matter.  We give several examples of non-standard scenarios that highlight the importance of achieving this complementarity through a variety of experiments with comparable sensitivity.   \\begin{figure}[t] \\includegraphics[width=3.25in]{fig1} \\caption{Solar WIMP capture rate as a function of WIMP mass for two values of the SD WIMP-proton scattering cross section.  Evaporation (see text) is not included. \\label{fig:caprate}} \\vspace{-1.\\baselineskip} \\end{figure} ",
        "conclusions": "We propose a new probe of the SD WIMP-nucleon scattering cross section, using low-energy neutrinos produced through pion multiplication and decay following WIMP annihilations in the Sun.  We estimate the prospects for Super-K, finding sensitivity to $\\sigma_{\\chi P}^{\\rm SD}$ in a range competitive to that of other, rather different, experiments.  Importantly, our results are nearly insensitive to the annihilation final states and the details of the astrophysical inputs (in standard scenarios~\\cite{Rott:2011fh}); non-standard scenarios can be different and are mentioned below.  In addition, the sensitivity easily extends to the region of low masses, which is currently of great interest and is challenging to probe with other methods.  A dedicated study by the Super-K collaboration, using more data, neutrino-electron scattering signals for all flavors, and full energy spectra and angular distributions, should give immediate improvements over our estimates.  \\begin{figure}[t] \\includegraphics[width=3.25in]{fig4} \\caption{New sensitivity to $\\sigma_{\\chi p}^{\\rm SD}$ using low-energy $\\bar{\\nu}_e$ in Super-K.  Possible improvements using neutron tagging with Gadolinium (details see text) are included in the projected sensitivity for 4~years of Hyper-Kamiokande~\\cite{Abe:2011ts} data.  For $m_\\chi \\lesssim 4$~GeV, evaporation of solar WIMPs will degrade the signal.  Published upper limits (90\\% C.L.) from direct scattering~\\cite{Archambault:2012pm, Kim:2012rz, Felizardo:2010mi} and high-energy neutrino searches are shown~\\cite{Tanaka:2011uf, Aartsen:2012kia}, along with the possible DAMA/LIBRA signal region~\\cite{Bernabei:2010mq, Savage:2008er}. \\label{fig:sens}} \\vspace{-1.\\baselineskip} \\end{figure}  In the future, Super-K may be enhanced with dissolved gadolinium to allow neutron detection and thereby better separation of signals and backgrounds~\\cite{Beacom:2003nk}.  If a combination of techniques removed the backgrounds, then the sensitivity to the neutrino signal and hence also $\\sigma_{\\chi P}^{\\rm SD}$ could improve by $\\sim 15$, beyond which even Super-K is too small to expect any signal events.  The quoted improvement factor is obtained with inverse beta decay alone assuming a signal event detection in a zero background environment. Hyper-Kamiokande~\\cite{Abe:2011ts}, which is intended to be about 25 times larger than Super-K and which also may have gadolinium, could potentially improve on present estimates by up to $\\sim 15 \\times 25 \\sim 375$. Figure~\\ref{fig:sens} shows our rough Hyper-Kamiokande sensitivity estimate under the assumption that backgrounds can be reduced and that only the inverse beta decay detection channel is utilized. Further studies on how to reduce backgrounds as well as contributions from other detection channels are needed but are beyond the scope of this paper. Proposed large liquid scintillator~\\cite{Wurm:2011zn} or liquid argon~\\cite{Rubbia:2009md} detectors would also have interesting sensitivity.  Direct, indirect, and collider probes of the SD WIMP-proton scattering cross section rely on different assumptions and hence are complementary.  Multiple methods with comparable sensitivity are needed to test results from one search against the others.  These tests could provide deep insights into the astrophysical distributions and particle properties of WIMPs.  If common assumptions about dark matter are incorrect, then the relative power of different methods could change dramatically.  A dark matter disk with low-velocity WIMPs could enhance both types of neutrino signals~\\cite{Bruch:2009rp}.  A time-varying WIMP flux from dark matter substructures could alter the relationship between neutrino and direct nuclear scattering signals~\\cite{Koushiappas:2009ee}.  WIMP annihilations through a new low-mass force carrier that decays only into low-mass SM particles~\\cite{ArkaniHamed:2008qn} could have vanishing high-energy but strong low-energy neutrino signals.  These are just a few examples.  Most generally, a combination of experiments will be required to reconstruct or constrain the complete couplings of the SM to dark matter, whether WIMPs or something more exotic.  \\medskip {\\bf Note added:} As this work was being completed, we learned of an independent study with similar results by Bernal, Martin-Albo, and Palomares-Ruiz~\\cite{Bernal:2012qh}. Both works were submitted simultaneously to arXiv."
    },
    "1208/1208.0016_arXiv.txt": {
        "abstract": "The super-massive black-holes in the centers of Active Galactic Nuclei (AGN) are surrounded by obscuring matter that can block the nuclear radiation. Depending on the amount of blocked radiation, the flux from the AGN can be too faint to be detected by currently flying hard X-ray (above 15 keV) missions. At these energies only $\\sim$1\\% of the intensity of the Cosmic X-ray Background (CXB) can be resolved into point-like sources that are AGNs. In this work we address the question of the undetected sources contributing to the CXB with a very sensitive and new hard X--ray survey: the {\\em SIX} survey that is obtained with the new approach of combining the {\\em Swift}/BAT and INTEGRAL/IBIS X--ray observations. We merge the observations of both missions. This enhances the exposure time and reduces systematic uncertainties. As a result we obtain a new survey over a wide sky area of 6200 deg$^{2}$ that is more sensitive than the surveys of  {\\em Swift}/BAT or INTEGRAL/IBIS alone. Our sample comprises 113 sources: 86 AGNs (Seyfert-like and blazars), 5 galaxies, 2 clusters of galaxies, 3 Galactic sources, 3 previously detected unidentified X-ray sources, and 14 unidentified sources. The scientific outcome from the study of the sample has been properly addressed to study the evolution of AGNs at redshift below 0.4. We do not find any evolution using the 1/V$_{max}$ method. Our sample of faint sources are suitable targets for the new generation hard X-ray telescopes with focusing techniques. ",
        "introduction": "In the view of the so-called AGN unified model \\citep{antonucci93,urry95} a super-massive black hole (SMBH) harbored at the center of the AGN powers the nuclear activity. The region where the activity takes place can be observed from different viewing angles. Therefore depending on the orientation of the AGN the observer's line of sight intercepts different amounts of the optically thick gas--dust structure (torus) that surrounds the SMBH. The nuclear radiation at optical/UV and X-ray wavelengths is efficiently absorbed by the torus. The amount of obscuring matter (N$_{H}$ column density associated to the torus) can be best inferred by X-ray spectra of the AGNs. X-ray surveys are therefore powerful tools for AGN population studies. The bias of X-ray surveys strongly depends on the column density associated to the sources and the survey sensitivity: the larger the column density and the worse the flux sensitivity, the better the low--absorbed sources are selected. Such selection effect is negligible for unabsorbed sources (exhibiting N$_{H}$ $<$ 10$^{22}$ cm$^{-2}$) while it affects the absorbed sources (exhibiting N$_{H}$ $>$ 10$^{22}$ cm$^{-2}$) and it is magnified for sources with column densities N$_{H}$ $\\ge$ 1.5 $\\times$ 10$^{24}$ cm$^{-2}$. This latter value corresponds to the inverse of the Thompson cross-section ($\\sigma_{T}^{-1}$) and the optical depth unity for Compton scattering. Absorbed sources affected by such high column densities are defined as \"Compton-thick\". This plays an important role in nowadays most sensitive AGN X-ray surveys that are performed by {\\em XMM-Newton} and {\\em Chandra} in the energy range $\\sim$0.5 - 10 keV \\citep{brandt01,alexander03,cappelluti09,xue11}. At these energies less than a mere 10\\% of the nuclear radiation is energetic enough to pierce through the absorbing Compton-thick torus \\citep{gilli07}. On the other hand the efficiently absorbed optical/UV radiation heats the dust of the obscuring medium, that is expected to waste the absorbed radiation in form of IR emission. Indeed, an IR--excess due to warm dust heated by obscured AGNs has been found \\citep{fadda02}. Infrared power-law selected samples in {\\em Chandra} Deep Fields are promising AGN--candidates \\citep{alonso-herrero06,donley07}. The drawback of the IR selection is that the majority of the detected sources are not AGNs. Furthermore this approach seems to sample best the sources within redshift 1--3 \\citep{donley07}. This is the same redshift range in which {\\em Chandra} and {\\em XMM-Newton} are preferentially selecting most AGNs in their deep surveys \\citep{brandt-hasinger05,hasinger08}. Instead the redshift space at z $<$ 0.4 is so far poorly explored despite extensive studies \\citep{markwardt05,beckmann06,sazonov07,ajello08c,tueller08,bird10,cusumano10}.  The low-redshift (z $<$ 0.4) Universe is best fathomed at hard X--ray energies ($>$ 15 keV). With the advent of the INTEGRAL \\citep{winkler03} and the {\\it Swift} \\citep{gehrels04} missions, the selection of local AGNs through their hard X--ray ($>$15 keV) emission has proven to be an extremely powerful technique over the last few years. INTEGRAL and {\\it Swift} carry coded-mask telescopes on board, namely the Imager on--Board the INTEGRAL Satellite \\citep[IBIS:][]{ubertini03} and the Burst Alert Telescope \\citep[BAT:][]{barthelmy05} respectively. IBIS has two detector layers. One of which is the INTEGRAL Soft Gamma-Ray Imager \\citep[ISGRI:][]{lebrun03}. IBIS/ISGRI and BAT have two major advantages: 1) they have a huge field of view, hence allowing to sample an adequate number of AGNs at low-redshift 2) they operate at energies above 15 keV, hence allowing detecting the photons having enough penetrating power to pierce efficiently even through the Compton-thick torus. A further and major advantage in sampling photons above 15 keV from AGNs comes from the emitting source itself. Indeed, a broad continuum bump, so-called \"Compton-reflection bump\", peaking at energies between 20 - 30 keV is produced by reflection of the primary nuclear radiation on the inner side of the obscuring gas \\citep{george91,gilli07}. This spectral component has been found to be dominant in nearby heavily  obscured AGNs \\citep{comastri07}. The Compton-reflection component also plays an important role in reproducing the shape and intensity of the CXB \\citep{rogers91,gilli01,ueda03}, that peaks at 30 keV \\citep[for most recent measurements see:][]{ajello08c,moretti09,tuerler10}. Estimates based on observations with PDS \\citep{frontera97} on board the {\\it Beppo}SAX \\citep{boella97} satellite predict that Compton-thick AGNs should dominate over unobscured AGNs in the local Universe \\citep{matt00}. This makes IBIS/ISGRI and BAT well suited instruments for detecting obscured AGNs in the local Universe. IBIS/ISGRI and BAT both represent a major improvement for the imaging of the sky above 15 keV. However coded-mask detectors suffer from heavy systematic effects (errors) preventing them from reaching their theoretical limiting sensitivity \\citep{skinner08}. Furthermore by design they block $\\sim$50\\% of the incident photons causing an increase of the statistical noise. These are the reasons that make the detection of extragalactic sources, that are mostly faint, still challenging to undertake. Here we describe an alternative approach which has been developed ad hoc to improve the sensitivity of extragalactic hard X--ray surveys by using IBIS/ISGRI and BAT.  In this paper we show that {\\it Swift}/BAT and INTEGRAL/IBIS observations can be merged to obtain a more sensitive survey that is able to sample limiting fluxes of $\\sim$3.3 $\\times$ 10$^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the 18 - 55 keV energy range. We call this the {\\em SIX} survey, that stands for {\\it Swift}--INTEGRAL hard X-ray survey. The {\\em SIX} survey extends over a wide sky area of 6200 deg$^{2}$ and it is used to obtain a small and persistent sample of faint sources. This enables the construction of the number density (log $N$--log $S$) as well as developing the X-ray luminosity function (XLF) for AGNs. In addition we estimate the contribution of this sample of AGNs to the intensity of the unresolved fraction of the CXB. Throughout this paper we adopt the cosmological parameters of: H$_{0}$ = 70 h$_{70}$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$ = 0.73. ",
        "conclusions": "A new approach has been developed to survey the sky at hard X-ray energies (18--55 keV energy band) by combining the observations of {\\it Swift}/BAT and INTEGRAL/IBIS resulting in the {\\em SIX} survey. First we have performed the independent surveys for both instruments. Then we have resampled, cross-calibrated and merged them. As a result of combining the observations from two different telescopes, statistical and systematic uncertainties caused by the high background level of their coded-mask detectors are minimized. In turn the {\\em SIX} survey is more sensitive, like a survey from a virtual new hard X-ray mission.\\\\ We applied the survey method to 6200 deg$^{2}$ of sky ($\\sim$ 20\\% of the entire extragalactic sky) sampling 113 sources that are: 74 Seyfert-like AGNs, 12 blazars, 5 galaxies, 2 clusters of galaxies, 3 Galactic sources, 3 previously detect X-ray sources, and 14 unidentified sources (of which 7 are newly detected without any counterpart and 7 are of uncertain association). No false detections due to statistical or systematic fluctuations are expected. The sources are identified through their soft X-ray counterparts and with {\\em Chandra} follow up observations (CXC AO-12). Unidentified sources are being followed-up in {\\em Chandra} AO-13. Among the AGN sample only two sources are Compton--thick, accounting for $\\sim$3\\% of the entire sample.\\\\ The number density of our identified sources is $\\sim$4 times greater than in our control sample of \\cite{ajello09b}. Even though this represents only a minor fraction of the CXB, the sensitivity improvement with respect to previous measurements is better than a factor of 2.\\\\ The fraction of absorbed AGN decreases with increasing luminosities. Although the redshift dependence is marginally significant, we find a mild decrease of the fraction of obscured AGNs with increasing redshift. These results require that the covering factor of the torus surrounding the SMBH changes at least with luminosity.\\\\ Only robustly identified AGNs were used in our XLF. The data are well represented by a double power-law model and do not show any evolution in density or in luminosity. The non-evolving XLF model fits our data best.\\\\ Based on our results we predict the number density of $\\sim$100 Seyfert--like AGNs that the upcoming NuSTAR mission can detect in 1 deg$^{2}$ of surveyed sky area at a limiting flux of $10^{-14}\\;erg\\;cm^{-2}\\;s^{-1}$.\\\\"
    },
    "1208/1208.4323.txt": {
        "abstract": "We report the first weak-lensing detection of a large-scale filament funneling matter onto the core of the massive galaxy cluster MACSJ0717.5+3745.  Our analysis is based on a mosaic of 18 multi-passband images obtained with the Advanced Camera for Surveys aboard the Hubble Space Telescope, covering an area of $\\sim 10 \\times 20$ arcmin$^{2}$. We use a weak-lensing pipeline developed for the COSMOS survey, modified for the analysis of galaxy clusters, to produce a weak-lensing catalogue. A mass map is then computed by applying a weak-gravitational-lensing multi-scale reconstruction technique designed to describe irregular mass distributions such as the one investigated here. We test the resulting mass map by comparing the mass distribution inferred for the cluster core with the one derived from strong-lensing constraints and find excellent agreement.  Our analysis detects the MACSJ0717.5+3745 filament within the 3~sigma detection contour of the lensing mass reconstruction, and underlines the importance of filaments for theoretical and numerical models of the mass distribution in the Cosmic Web. We measure the filament's projected length as $\\sim$ 4.5 $h_{74}^{-1}$ Mpc, and its mean density as $(2.92 \\pm 0.66)\\times10^{8}~h_{74}$ M$_{\\odot}$ kpc$^{-2}$. Combined with the redshift distribution of galaxies obtained after an extensive spectroscopic follow-up in the area, we can rule out any projection effect resulting from the chance alignment on the sky of unrelated galaxy group-scale structures. Assuming plausible constraints concerning the structure's geometry based on its galaxy velocity field, we construct a 3D model of the large-scale filament. Within this framework, we derive the three-dimensional length of the filament to be 18~$h_{74}^{-1}$ Mpc.  The filament's deprojected density in terms of the critical density of the Universe is measured as $(206 \\pm 46)\\times \\rho_{\\rm crit}$, a value that lies at the very high end of the range predicted by numerical simulations. Finally, we study the distribution of stellar mass in the field of MACSJ0717.5+3749 and, adopting a mean mass-to-light ratio $\\langle M_{\\ast}/L_{K}\\rangle$ of $0.73 \\pm 0.22$ and assuming a Chabrier Initial-Mass Function, measure a stellar mass fraction along the filament of $(0.9 \\pm 0.2)$\\%, consistent with previous measurements in the vicinity of massive clusters. ",
        "introduction": "In a Universe dominated by Cold Dark Matter (CDM), such as the one parameterised by the $\\Lambda$CDM concordance cosmology, hierarchical structure formation causes massive galaxy clusters to form through a series of successive mergers of smaller clusters and groups of galaxies, as well as through continuous accretion of  surrounding matter. % channeled through filaments. N-body simulations of the dark-matter distribution on very large scales \\citep{bond96,YS96,aragoncalvo07,hahn07} predict that these processes of merging and accretion occur along preferred directions, i.e., highly anisotropically. The result is the ``cosmic web'' \\citep{bond96}, a spatially highly correlated structure of interconnected filaments and vertices marked by massive galaxy clusters. Abundant observational support for this picture has been provided by large-scale galaxy redshift surveys \\citep[e.g.,][] {GH89, york00,colless01} showing voids surrounded and connected by filaments and sheets of galaxies.  A variety of methods have been developed to detect filaments in surveys, among them a ``friends of friends'' algorithm \\citep[FOF,][]{HG82} combined with ``Shapefinders'' statistics \\citep{SJ03}; the ``Skeleton'' algorithm \\citep{novikov06,sousbie06a}; a two-dimensional technique developed by \\citet{moody83}; and the Smoothed Hessian Major Axis Filament Finder \\citep[SHMAFF,][]{bond10}.  Although ubiquitous in large-scale galaxy surveys, filaments have proven hard to characterise physically, owing to their low density and the fact that the best observational candidates often turn out to be not primordial in nature but the result of recent cluster mergers. Specifically, attempts to study the warm-hot intergalactic medium \\citep[WHIM,][]{CO99}, resulting from the expected gravitational heating of the intergalactic medium in filaments, remain largely inconclusive because it is hard to ascertain for filaments near cluster whether spectral X-ray features originate from the filament or from past or ongoing clusters mergers \\citep{kaastra06,rasmussen07,galeazzi09,williams10}. Some detections appear robust as they have been repeatedly confirmed \\citep{fang02,fang07,williams07} but are based on just one X-ray line. An alternative observational method is based on a search for filamentary overdensities of galaxies relative to the background \\citep{PD04,ebeling04}. When conducted in 3D, i.e., including spectroscopic galaxy redshifts, this method is well suited to detecting filament candidates. It does, however, not allow the determination of key physical properties unless it is supplemented by follow-up studies targeting the presumed WHIM and dark matter which are expected to constitute the vast majority of the mass of large-scale filaments. By contrast, weak gravitational lensing offers the tantalising possibility of detecting directly the total mass content of filaments \\citep{mead10}, since the weak-lensing signal arises from luminous and dark matter alike, regardless of its dynamical state.  %______________________________________________________ % Fig 1 %______________________________________________________ \\begin{figure}%[!h] \\includegraphics[width=84mm]{fig1.eps} \\caption{Area of our spectroscopic survey of MACSJ0717.5+3745. Outlined in red is the region covered by our Keck/DEIMOS masks; outlined in blue is the area observed with HST/ACS. Small circles correspond to objects for which redshifts were obtained; large filled circles mark cluster members. The black contours show the projected galaxy density (see Ma et al.\\ 2008).} \\label{fig:area0717} \\end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  Previous weak gravitational lensing studies of binary clusters found tentative evidence of filaments, but did not result in clear detections. One of the first efforts was made by \\citet{clowe98} who reported the detection of a filament apparently extending from the distant cluster RX\\,J1716+67 ($z=0.81$), using images obtained with the Keck 10m telescope and University of Hawaii (UH) 2.2m telescope. This filamentary structure would relate two distinct sub clusters detected on the mass and light maps. The detection was not confirmed though. Almost at the same time \\citet{kaiser98} conducted a weak lensing study of the supercluster MS0302+17 with the UH8K CCD camera on the Canada France Hawaii Telescope (CFHT). Their claimed detection of a filament in the field was, however, questioned on the grounds that the putative filament overlapped with both a foreground structure as well as with gaps between CCD chips. Indeed, \\citet{gavazzi04} showed the detection to have been spurious by means of a second study of MS0302+17 using the CFHT12K camera. A weak gravitational lensing analysis with MPG/ESO Wide Field Imager conducted by  \\citet{gray02}  claimed the detection of a filament in the triple cluster A901/902. The candidate filament appeared to connect two of the clusters and was detected in both the galaxy distribution and in the weak-lensing mass map. However, this detection too was of low significance and coincided partly with a gap between two chips of the camera. As in the case of MS0302+17, a re-analysis of the A901/A902 complex using high-quality HST/ACS images by \\citet{heymans08} failed to detect the filament and led the authors to conclude that the earlier detection was caused by residual PSF systematics and limitation of the KS93 mass reconstruction used in the study by \\citet{gray02}. A further detection of a filament candidate was reported by \\citet{dietrich05} based on a weak gravitational lensing analysis of the close double cluster A222/A223. However, as in other similar cases, the proximity of the two clusters connected by the putative filament raises the possibility of the latter being a merger remnant rather than primordial in nature.  %__________________________________________________________________ % TABLE 1 %__________________________________________________________________ \\begin{table*}%[!t] \\caption{Overview of the HST/ACS observations of MACSJ0717.5+3745. *: Cluster core; observed through F555W, rather than F606W filter.} \\label{tab:acs_obs} \\begin{center} \\begin{tabular}[h]{|ccccccc} \\hline\\\\[-5pt] &   & \\multicolumn{2}{c}{F606W} & \\multicolumn{2}{c}{F814W} &  \\\\ \\cline{3-4} \\cline{5-6} \\\\[-5pt] R.A.\\ (J2000) & Dec (J2000) & Date & Exposure Time (s) & Date & Exposure Time (s) & Programme  ID \\\\ \\hline\\\\[-5pt] 07 17 32.93 & +37 45 05.4 & \\,\\,\\,2004-04-02* & \\,\\,\\,4470* & 2004-04-02 & 4560 & \\,\\,\\,\\,9722 \\\\ 07 17 31.81 & +37 49 20.6 & 2005-02-08 & 1980 & 2005-02-08 & 4020 & 10420 \\\\ 07 17 20.38 & +37 47 07.5 & 2005-01-27 & 1980 & 2005-01-27 & 4020 & 10420 \\\\ 07 17 08.95 & +37 44 54.3 & 2005-01-27 & 1980 & 2005-01-27 & 4020 & 10420 \\\\ 07 17 43.23 & +37 47 03.1 & 2005-01-30 & 1980 & 2005-01-30 & 4020 & 10420 \\\\ 07 17 20.18 & +37 42 38.8 & 2005-02-01 & 1980 & 2005-02-01 & 4020 & 10420 \\\\ 07 17 54.26 & +37 44 49.3 & 2005-01-27 & 1980 & 2005-01-27 & 4020 & 10420 \\\\ 07 17 42.82 & +37 42 36.3 & 2005-01-24 & 1980 & 2005-01-25 & 4020 & 10420 \\\\ 07 17 31.39 & +37 40 23.3 & 2005-02-01 & 1980 & 2005-02-01 & 4020 & 10420 \\\\ 07 18 05.46 & +37 42 33.6 & 2005-02-04 & 1980 & 2005-02-04 & 4020 & 10420 \\\\ 07 17 54.02 & +37 40 20.6 & 2005-02-04 & 1980 & 2005-02-04 & 4020 & 10420 \\\\ 07 17 42.79 & +37 38 05.7 & 2005-02-05 & 1980 & 2005-02-05 & 4020 & 10420 \\\\ 07 18 16.65 & +37 40 17.7 & 2005-02-05 & 1980 & 2005-02-05 & 4020 & 10420 \\\\ 07 18 05.22 & +37 38 04.9 & 2005-02-05 & 1980 & 2005-02-05 & 4020 & 10420 \\\\ 07 17 53.79 & +37 35 52.0 & 2005-02-05 & 1980 & 2005-02-05 & 4020 & 10420 \\\\ 07 18 27.84 & +37 38 01.9 & 2005-02-08 & 1980 & 2005-02-08 & 4020 & 10420 \\\\ 07 18 16.40 & +37 35 49.1 & 2005-02-08 & 1980 & 2005-02-08 & 4020 & 10420 \\\\ 07 18 04.97 & +37 33 36.2 & 2005-02-09 & 1980 & 2005-02-09 & 4020 & 10420 \\\\ \\hline\\\\[-5pt] \\end{tabular}\\\\ \\end{center} \\end{table*} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  In this paper we describe the first weak gravitational analysis of the very massive cluster MACSJ0717.5+3745 \\citep[$z=0.55$;][]{edge03,ebeling04,ebeling07,ma08,ma09}. Optical and X-ray analyses of the system \\citep{ebeling04, ma08,ma09} find compelling evidence of a filamentary structure extending toward the South-East of the cluster core. Using weak-lensing data to reconstruct the mass distribution in and around MACSJ0717.5+3745, we directly detect the reported filamentary structure in the field of MACSJ0717.5+3745.  The paper is organized as follows. After an overview of the observational data in Section 2, we discuss the gravitational lensing data in hand in Section 3. The modeling of the mass using a multi-scale approach is described in Section 4. Results are discussed in Section 5, and we present our conclusions in Section 6. \\\\ \\\\ All our results use the $\\Lambda$CDM concordance cosmology with $\\Omega_{\\rm M}$ = 0.3, $\\Omega_{\\Lambda}$ = 0.7, and a Hubble constant $H_{0}$ = 74 km s$^{-1}$ Mpc$^{-1}$, hence 1\" corresponds to 6.065~kpc at the redshift of the cluster. Magnitudes are quoted in the AB system.  %______________________________________________________ % SECTION 2 :  OBSERVATIONS %______________________________________________________ ",
        "conclusions": "We present the results of a weak-gravitational lensing analysis of the massive galaxy cluster MACSJ0717.5+3745 and its large-scale filament, based on a mosaic composed of 18 HST/ACS images covering an area of approximately $10\\times 20$ arcmin$^2$. Our mass reconstruction method uses RRG shape measurements \\citep{rhodes07}; a multi-scale adaptive grid designed to follow the structures' K-band light and including galaxy-size potentials to account for cluster members; and the LENSTOOL software package, improved by the implementation of a Bayesian MCMC optimisation method that allows the propagation of measurement uncertainties into errors on the filament mass. As a critical step of the analysis, we use spectroscopic and photometric redshifts, as well as colour-colour cuts, all based on groundbased observations, to eliminate foreground galaxies and cluster members, thereby reducing dilution of the shear signal from unlensed galaxies.  A simple convergence map of the study area, obtained with the inversion method of \\cite{seitz95}, already allows the detection of the cluster core (at more than $6\\sigma$ significance) and of two extended mass concentrations (at 2 to 3$\\sigma$ significance) near the beginning and (apparent) end of the filament.  The fully optimised weak-lensing mass model yields the  surface mass density shown in Fig.~\\ref{fig:massmap_WL}. Its validity is confirmed by the excellent agreement between the mass of the cluster core measured by us, $M_{\\rm WL} (R<500 \\,{\\rm kpc}) = (1.04 \\pm 0.08)\\times 10^{15} \\,h_{74}^{-1}$ M$_{\\odot}$, and the one obtained in a strong-lensing analysis by \\cite{limousin11}, $M_{\\rm SL} (R<500 \\,{\\rm kpc}) = (1.06 \\pm 0.03)\\times 10^{15}\\,h_{74}^{-1}$ M$_{\\odot}$.  Based on our weak-lensing mass reconstruction, we report the first unambiguous detection of a large-scale filament fueling the growth of a massive galaxy cluster at a node of the Cosmic Web. The projected length of the filament is approximately $4.5\\,h_{74}^{-1}$ Mpc, and its mean mass surface density  $(2.92 \\pm 0.66) \\times 10^{8}\\, h_{74}$  M$_{\\odot}$ kpc$^{-2}$. We measure the width and mass surface density of the filament as a function of distance from the cluster centre, and find both to decrease, albeit with local variations due to at least four embedded mild mass concentrations. One of these, at a projected distance of 1.85\\,$h_{74}^{-1}$ Mpc from the cluster centre, coincides with an X-ray detected galaxy group. The filament is found to narrow at a cluster-centric distance of approximately $3.6\\,h_{74}^{-1}$ Mpc as it curves south and, most likely, also recedes from us. We find the cluster's mass surface density to decrease as $r^{-2}$ until the onset of the filament flattens the profile dramatically.  Following the analysis by Ebeling et al.\\ (2012, in preparation) we adopt an average inclination angle with respect to the plane of the sky of 75$^\\circ$ for the majority of the filament's length. For this scenario and under the simplifying assumption of a cylindrical cross-section we obtain estimates of the filament's three-dimensional length and average mass density in units of the Universe's critical density of $18\\,h^{-1}_{74}$ Mpc and $(206 \\pm 46)\\,\\rho_{\\rm crit}$, respectively. However,  additional systematic uncertainties enter since either quantity is sensitive to the adopted inclination angle. These values lie at the high end of the range predicted from numerical simulations \\citep[e.g.,][]{colberg05}, which is not unexpected given the extreme mass of MACSJ0717.5+3745.  The lensing surface mass density and the spectroscopic redshift distribution suggest the  consistent picture of an elongated structure at the redshift of the cluster. The galaxy distribution along the filament appears to be homogeneous, and at the cluster redshift \\citep[see][]{ebeling04}. This motivates our conclusion of an unambiguous detection of a large scale filament, and not a superposition of galaxy groups projected on the plane of the sky.  Finally, we measure the stellar mass fraction in the entire MACSJ0717.5+3745 field, using as a proxy the K-band luminosity of galaxies with redshifts consistent with that of the cluster-filament complex. We find $f_{\\ast} = (1.3 \\pm 0.4)\\%$ and $f_{\\ast} = (0.9 \\pm 0.2) \\%$ for a Salpeter and Chabrier IMF, respectively, in good agreement with previous results in the fields of massive clusters \\citep{leauthaud11}.  Our results show that, if shear dilution by unlensed galaxies can be efficiently suppressed, weak-lensing studies of massive clusters are capable of detecting and mapping the complex mass distribution at the vertices of the cosmic web. Confirming results of numerical simulations \\citep[e.g.][]{colberg99, colberg05,OH11}, our weak-lensing mass reconstruction shows that the contribution from large-scale filaments can be significant and needs to be taken into account in the modelling of mass density profiles. Expanding this kind of investigation to HST/ACS-based weak-lensing studies of other MACS clusters will allow us to constrain the properties of large-scale filaments and the dynamics of cluster growth on a sound statistical basis.  %______________________________________________________ % ACKNOLEDGMENTS %______________________________________________________"
    },
    "1208/1208.5069_arXiv.txt": {
        "abstract": "Type Ia supernovae (\\sneia) originate from the thermonuclear explosions of carbon-oxygen (C-O) white dwarfs (WDs). The single-degenerate scenario is a well-explored model of \\sneia\\ where unstable thermonuclear burning initiates in an accreting, Chandrasekhar-mass WD and forms an advancing flame. By several proposed physical processes the rising, burning material triggers a detonation, which subsequently consumes and unbinds the WD. However, if a detonation is not triggered and the deflagration is too weak to unbind the star, a completely different scenario unfolds. We explore the failure of the Gravitationally-Confined Detonation (GCD) mechanism of \\sneia, and demonstrate through 2D and 3D simulations the properties of failed-detonation SNe. We show that failed-detonation SNe expel a few 0.1 \\msolar\\ of burned and partially-burned material and that a fraction of the material falls back onto the WD, polluting the remnant WD with intermediate-mass and iron-group elements, that likely segregate to the core forming an WD whose core is iron rich. The remaining material is asymmetrically ejected at velocities comparable to the escape velocity from the WD, and in response, the WD is kicked to velocities of a few hundred km s$^{-1}$. These kicks may unbind the binary and eject a runaway/hyper-velocity WD. Although the energy and ejected mass of the failed-detonation SN are a fraction of typical thermonuclear SNe, they are likely to appear as sub-luminous low-velocity \\sneia. Such failed detonations might therefore explain or are related to the observed branch of peculiar \\sneia, such as the family of low-velocity sub-luminous SNe (SN 2002cx/SN 2008ha-like SNe). ",
        "introduction": "Type Ia supernovae (\\sneia) are among the most energetic explosions in the known universe, releasing $\\sim 10^{51}$ ergs of kinetic energy in their ejecta, and synthesizing $\\sim 0.7$ \\msolar\\ of radioactive \\iso{Ni}{56}. The discovery of the Phillips relation \\citep{pskovskii77, phillips93} enabled the use of \\sneia\\ as standardizable cosmological candles, and has ushered in a new era of astronomy leading to the discovery of the acceleration of the universe \\citep{1998AJ....116.1009R, schmidtetal98, 1999ApJ...517..565P}, and to the 2011 Nobel Prize in physics.  Models of normal \\sneia, such as the single degenerate (SD) model, focus on exploding the WD in order to produce the explosion energies, luminosities, and typical velocities observed in normal \\sneia. This is accomplished either by consuming enough of the WD with the initial subsonic buring phase --- or deflagration phase --- to unbind the WD as theorized by the Pure Deflagration (PD) \\citep{2005ApJ...623..337G, 2005A&A...431..635R} model, or consuming the entire WD by a detonation triggered by the deflagration phase as posited by the deflagration-to-detonation transition (DDT) \\citep{1991A&A...245..114K, 2004PhRvL..92u1102G, 2005ApJ...623..337G}, the Pulsating Reverse Detonation (PRD) \\citep{2009ApJ...695.1244B, 2009ApJ...695.1257B}, and the GCD \\citep{2004astro.ph..5162C,2008ApJ...681.1448J, 2009ApJ...693.1188M, 2012ApJ...759...53J} models of \\sneia.  In the following, we present a novel variant of SD model of \\sneia\\ in which the deflagration is too weak to unbind the star\\footnote{We acknowledge related work by \\citet{2012arXiv1210.5243K} that appeared as this article was going to press.} and a detonation is not triggered by any of the proposed mechansisms, resulting in the survival of a bound remnant of the original WD. We present for the first time predictions of these failed-detonation (FD) SNe from 2D and 3D simulations. We show that FD models have numerous remarkable implications for the observable properties of the resulting explosion and its outcomes. These include the production of a family of peculiar \\sneia\\ events with low expansion velocities, low luminosities and low ejecta-mass --- whose properties are broadly consistent with the observed properties of a branch of peculiar \\sneia\\ similar to SN 2002cx and/or SN 2008ha. Even more remarkably, the remnant WD receives a large velocity kick from the asymmetric nature of the deflagration, and is enriched with both intermediate-mass (IME's) and iron-group elements (IGE's), forming a peculiar WD with a heavey/iron-rich core.  Previous work has suggested that though the PD model has shortcomings explaining normal \\sneia, they may explain 2002cx-like SNe \\citep{branchetal04}. The WD, however, is fully incinerated in these models, producing a Chandrasekhar-mass of ejecta. Such models might therefore not be able to explain the large diversity recently observed among SNe of this peculiar class of SNe. Additionally, a study by \\citet{livneetal05} of initial conditions for the PD model produced situations that they termed ``fizzles'' which did not produce a healthy PD explosion and left the WD bound. They did not pursue these models beyond the scope of their study, nor did they relate these fizzles to peculiar \\sneia. ",
        "conclusions": "} \\subsection{Sub-luminous Low-Velocity SNe}  The most prominent features of the FD models are the low mass and low velocity of the ejecta, which translate into the production of typically sub-luminous, low-velocity \\sneia. It is therefore natural to examine whether SNe with such characteristics have already been discovered. In particular, one may explore peculiar SNe exhibiting either extremely low mass ejecta, such as SN 2008ha \\citep{fol+09,val+09}, or low ejecta velocity such as SN 2002cx-like SNe \\citep{lietal03, branchetal04, jhaetal06}.  Normal \\sneia\\ differ in their ejecta velocities as measured in some standard method, but they generally fall between $9000-14000$ km s$^{-1}$ near peak luminosity, with similar dispersion at later times (as derived from the Si II line \\citep{ben+05}). The velocities of even the lowest velocity SNe in the \\citet{ben+05} sample much exceed the mass-averaged FD velocities listed in table \\ref{tab:sims}. We also note a trend in our simulations of more energetic and likely more luminous (larger IGE yield) SNe to be accompanied by higher ejecta-velocities over almost an order of magnitude in kinetic energy. The only other type of SNe with such low expansion velocities are the branch of peculiar type Ia SNe, named for the prototype for this class of supernovae, 2002cx \\citep{lietal03, branchetal04, jhaetal06}; such SNe may also have an energy-velocity correlation \\citep{mcc+10} as observed in our simulations.  SN 2002cx-like events are characterized by luminosities which lie too low in comparison to the Phillips relation for Branch-normal Ia events \\citep{lietal03}, low photospheric velocities \\citep{lietal03}, weak intermediate-mass element lines \\citep{branchetal04}, and late-time optical nebular spectra dominated by narrow Fe II lines \\citep{branchetal04, jhaetal06}. Since the discovery of SN 2002cx, a number of other 2002cx-like events have been discovered, including 2002es, 2005P, 2005hk \\citep{chornocketal06}, 2008ge \\citep{foleyetal10} and 2008ha \\citep{fol+09,val+09}. The latter event (SN 2008ha), in particular, is consistent with extremely low mass ejecta and energetics. We predict the FD models to produce similar properties to those characterizing SN 2002cx like SNe,  given the low expansion velocities and the low estimated \\iso{Ni}{56} yield, and potentially even explaining SN 2008ha like events with low mass ejecta.  Though our initial set of simulations is limited, the robust features of FD's, including low velocity ejecta, the expected low luminosity (due to the small yield of \\iso{Ni}{56}) and their low mass ejecta (comparable to that SN 2008ha) make them tantalizing candidate progenitors for this branch of peculiar SNe. Note that the single degenerate origin of such SNe is also consistent with the overall typically young (but not \\emph{necessarily} young; \\citeauthor{fol+10} 2010) environments found for SN-2002cx like SNe, compared to the expectations from, e.g. core-collapse SNe (only very young environments).  \\subsection{WD's with Heavy/Iron-Rich Cores\\label{subsec:fecores}}  In our FD scenario, a large amount of burnt material falls back to the remnant WD. From table \\ref{tab:sims}, the WD may incorporate as much as $0.3$ $M_{\\odot}$ of IGE's and $0.07$ $M_{\\odot}$ of IME's of fallback material, together comprising as much as $\\sim 18\\%$ of the remnant C-O WD. In time, these elements are likely to gravitationally settle to the WD core, creating WD's with iron-rich/heavy core. The existence of iron-core WD's has been considered before, with even the potential observation of such WDs \\citep{pro+98,cat+08}. The FD scenario therefore provides a novel evolutionary scenario for the formation of these iron/heavy-core C-O WD's. A somewhat related scenario of failed SN was suggested for the formation of O-Ne-Mg WDs with iron cores \\citep{ise+91}.  \\subsection{WD Natal Kicks}  FD's produce a highly asymmetric ejection of material. This is not unique amongst various models for SNe explosions. However, in our FD case, the WD survives the explosion. Considering momentum conservation, this gives rise to a unique outcome, namely that the surviving WD is kicked at very high velocities, ranging hundreds of km s$^{-1}$. The FD scenario suggests the existence of strong WD natal kicks, and provides an interesting prediction per the existence of hypervelocity WDs. Taken together, the potential existence of a a heavy core WD (discussed in section \\ref{subsec:fecores}), and the high ejection velocity produce a highly peculiar object, which, if observed may provide a possibly unique smoking gun signature. One should note, however, that the population of halo WDs may also have relatively high velocities, and it might therefore be difficult to pinpoint the kinematic property as related to a natal kick (unless the WD is massive and young; an unlikely possibility for the old population of halo WDs).  We note that velocities of hundreds of km s$^{-1}$ could be larger than the orbital velocities of the SN binary progenitor, and a kick velocity of such magnitude can therefore unbind the binary. Various binary configurations have been explored for the single degenerate progenitor models, including progenitors with MS and RG companions \\citep{1996ApJ...470L..97H, 2000ApJS..128..615M}. We conclude that the range of WD kick velocities could either unbind the binary (more likely for WD-MS binaries), or leave behind a bound WD binary (more likely for the WD-RG binaries). The latter case could lead to the formation of a very compact, but potentially eccentric WD-binary, which would be difficult to produce through other channels of binary evolution."
    },
    "1208/1208.1821_arXiv.txt": {
        "abstract": "Using hydrodynamic simulations, we investigate the physical properties of gaseous substructures in barred galaxies and their relationships with the bar strength. The gaseous medium is assumed to be isothermal and unmagnetized. The bar potential is modeled as a Ferrers prolate with index $n$. To explore situations with differing bar strength, we vary the bar mass $\\fbar$ relative to the spheroidal component as well as its aspect ratio $\\R$. We derive expressions as functions of $\\fbar$ and $\\R$ for the bar strength $\\Qb$ and the radius $\\RQ$ where the maximum bar torque occurs. When applied to observations, these expressions suggest that bars in real galaxies are most likely to have $\\fbar\\sim0.25$--$0.50$ and $n\\simlt1$. Dust lanes approximately follow one of $\\xone$-orbits and tend to be more straight under a stronger and more elongated bar, but are insensitive to the presence of self-gravity.  A nuclear ring of a conventional $\\xtwo$ type forms only when the bar is not so massive or elongated. The radius of an $\\xtwo$-type ring is generally smaller than the inner Lindblad resonance, decreases systematically  with increasing $\\Qb$, and slightly larger when self-gravity is included. This evidences that the ring position is not determined by the resonance but by the amount of angular momentum loss at dust-lane shocks. Nuclear spirals exist only when the ring is of the $\\xtwo$-type and sufficiently large in size. Unlike the other features, nuclear spirals are transient in that they start out as being tightly-wound and weak, and then due to the nonlinear effect unwind and become stronger until turning into shocks, with an unwinding rate higher for larger $\\Qb$. The mass inflow rate to the galaxy center is found to be less than $0.01\\Aunit$ for models with $\\Qb\\simlt0.2$, while becoming larger than $0.1\\Aunit$ when $\\Qb\\simgt0.2$ and self-gravity is included. ",
        "introduction": "}  The inner parts of barred galaxies contain interesting gaseous substructures that include dust lanes, nuclear rings, and nuclear spirals (e.g., \\citealt{san76,sak99,kna02, mar03a,mar03b,she05,pri05,mar06,com09,com10,maz11,hsi11}). These substructures are thought to form as a result of gas redistribution in galaxies initiated by non-axisymmetric bar torque (e.g., \\citealt{com85,but86,shl90,gar91,but96,com01}). Understanding their formation and evolution is therefore of crucial importance to understand how the interstellar gas in barred galaxies is driven inward to affect star formation activities in the nuclear regions and potentially fueling of active galactic nuclei (AGN). Since their spatial locations and shapes are likely to be determined by the bar strength as well as the underlying rotation curve (e.g., \\citealt{ath92b,pee06,com09,com10}), they may also be used to probe the mass distributions in barred galaxies that are not directly observable.   While the bar substructures have been observed for a long time (e.g., \\citealt{pea17,san61}), it is during recent years that high-quality data on their physical properties have been compiled and compared with bar characteristics (e.g., \\citealt{kna02,pee06,maz08,com09,com10,maz11}). There is increasing observational evidence that the shape and size of the bar substructures are determined primarily by the bar strength $\\Qb$. For instance, \\citet{kna02} measured the curvature angles $\\Dalp$ of dust lanes for a sample of 9 barred-spiral galaxies, and found that stronger bars favor more straight dust lanes, confirming a theoretical prediction of \\citet{ath92b}.  By extending the sample size to 55 galaxies, \\citet{com09} confirmed the results of \\citet{kna02}, although they also found that a large scatter in the fit of the $\\Dalp$--$\\Qb$ relationship can be reduced if the bar ellipticity is also considered in the fit.  More recently, \\citet{com10} measured the sizes and shapes of nuclear rings in a sample of 107 galaxies, finding that `stronger bars host smaller rings' and that the ring ellipticity is in the range of 0--0.4. \\citet{maz11} reported from the analysis of 13 nuclear rings that the ring size is well correlated with the compactness of the galaxy mass distribution such that higher mass concentration implies a smaller ring. On the other hand, by analyzing dust morphology in the nuclear regions of 75 galaxies, \\citet{pee06} found that tightly-would nuclear spirals are preferentially found in galaxies with a weak bar.  On the theoretical side, the formation and evolution of bar substructure has been extensively studied using numerical simulations on hydrodynamical models (e.g. \\citealt{san76,rob79,van81,ath92b,pin95,eng97,pat00,mac02,mac04a, mac04b,reg03,reg04,ann05,lin08,tha09}) and magnetized models (e.g., \\citealt{kul09,kul10,kul11,ks12}). In particular, \\citet{ath92b} confirmed the early notion of \\citet{pre62} that dust lanes are shocks in the gas flows. She also showed that dust lanes are more straight when the bar potential is stronger.  Due to lack of numerical resolution and/or large numerical viscosity, however, early numerical works published before the middle of 1990s were unable to capture the formation of nuclear rings and spirals accurately. \\citet{pin95} was the first who focused on the ring formation using a grid-based code with negligible numerical viscosity. Unfortunately, however, their numerical results were, at least quantitatively, contaminated by a trivial sign error in the evaluation of the bar forces, as identified by \\citet{kim12}. More recently, \\citet{ann05} and \\citet{tha09} used SPH simulations to study gas dynamics in the central regions of barred galaxies, but they were unable to resolve nuclear spirals because most particles in their simulations are captured into a ring, leaving only a small number of particles inside the ring.  Very recently, \\citet[hereafter Paper I]{kim12} corrected the error in \\citet{pin95} and revisited the issue of the bar-substructure formation using high-resolution simulations by varying the gas sound speed $\\cs$ and the mass of a central black hole (BH). Paper I found that the corrected bar force naturally allows the development of nuclear spirals inside a ring, which was unseen in \\citet{pin95}. In addition, Paper I showed that the ring size and the mass inflow rate $\\Mdot$ toward the galaxy center depend rather sensitively on $\\cs$. In models with small $\\cs$, the rings are narrow and located away from the center. This prevents further inflows of the gas to the central regions, making $\\Mdot$ smaller than $10^{-4}\\Aunit$. On the the other hand, models with large $\\cs$ have a small and broad ring, resulting in $\\Mdot\\simgt10^{-2}\\Aunit$. Paper I further confirmed that the prediction of \\citet{but96} that the shape of nuclear spirals depend on the sign of the $d(\\Omega-\\kappa/2)/dr$ curve such that they are trailing (leading) where $d(\\Omega-\\kappa/2)/dr$ is negative (positive). This suggests that nuclear spirals, if exist, are likely to be trailing in galaxies with supermassive BHs.  Since the models studied in Paper I were restricted to a particular set of the bar parameters, namely, with the bar mass $\\sim30\\%$ of the spheroidal component (i.e., bulge plus bar) and the aspect ratio $\\R=2.5$, they were unable to study the variations in the size and shape of bar substructures with differing bar strength.  In addition, the models in Paper I are all non-self-gravitating, so that the effects of self-gravity on bar substructures have yet to be explored. Although \\citet{ath92b} considered diverse bar models with different aspect ratio and bar quadrupole moment $\\Qm$, her models were unable to resolve nuclear rings and spirals due to insufficient resolution. While \\citet{reg03,reg04} also ran a large number of numerical simulations with varying $\\Qm$ and $\\R$, their numerical results were compromised by the incorrect bar forces of \\citet{pin95}, as mentioned above.  In this paper, we present the results of a series of hydrodynamic simulations to investigate the properties of bar substructures that form. This work extends Paper I in two ways: (1) by considering various bar models with differing mass and aspect ratio and (2) by including gaseous self-gravity, while fixing the sound speed and the BH mass. We measure the curvatures of dust lanes, sizes of nuclear ring, shapes of nuclear spirals, and mass inflow rates in simulations, and study their relationships with the bar parameters. Our main objective is to explore the parametric dependence of the properties of bar substructures on the bar strength and self-gravity. This will allow us to provide physical explanations for the formation of bar substructures especially for nuclear rings and spirals, which was previously uncertain. We also compare our numerical results with observations of barred-spiral galaxies that are currently available.  Our models with varying bar strength are useful to address an important question as to what determines the position of nuclear rings. Observations indicates that nuclear rings are typically located near the inner Lindblad resonance (ILR) when there is a single ILR, or between the ILRs when there are two ILRs (e.g., \\citealt{com85, kna95,com10}). This has often been interpreted as an indication that nuclear rings form as a consequence of the resonant interactions of the gas with the ILRs (\\citealt{com96,but96}; see also \\citealt{reg03}). This idea was motivated by the notion that the bar torque changes its sign whenever crossing each ILR such that in the case of a single ILR, for example,  the gas inside (outside) the ILR receives a positive (negative) torque and thus moves outward (inward). This idea of the resonance-driven ring formation requires that the bar torque dominates thermal and ram pressures of the gas in governing the gas motions both inside and outside the ILR.  It is clear that the bar torque dominates outside the ILR, inducing dust-lane shocks and initiating radial gas inflows. However, it is uncertain if that is also the case inside the ILR.  The bar potential becomes increasingly axisymmetric toward the galaxy center, while the inflowing gas usually has large momentum and is thus unlikely to stall at the ILR. By measuring the ring positions and comparing them with the ILR radii in our models, we directly test whether the ILR is really responsible for the formation of nuclear rings.  This paper is organized as follows: In Section 2, we describe our model parameters and numerical methods.  In Section 3, we define and evaluate the bar strength $\\Qb$ of our galaxy models as well as the radius $\\RQ$ where the maximum bar torque occurs.  We provide algebraic expressions for $\\Qb$ and $\\RQ$ for various bar models. In Section 4, we quantify the properties of bar substructures and explore their correlations with the bar strength. We also measure the mass inflows rates. Finally, in Section 5, we summarize our results and discuss them in comparison with observations. ",
        "conclusions": "\\label{sec:sum}  We have presented the results of high-resolution hydrodynamic simulations on the formation and evolution of gaseous substructures in barred galaxies with varying bar strength. We initially consider an infinitesimally-thin, isothermal, unmagnetized gas disk with uniform surface density embedded in the external gravitational potential. We run both non-self-gravitating and self-gravitating models, but the effects of star formation and feedback are not considered in the present work.  In order to focus on the effects of the bar parameters, we fix the gas sound speed to $\\cs=10\\kms$ and the BH mass to $\\MBH=4\\times10^7\\Msun$ that affects the rotation curve near the galaxy center, and vary two parameters: the bar mass measured by $\\fbar$ (Eq.\\ [\\ref{eq:fbar}]) relative to the spheroidal component and its aspect ratio $\\R$ (see Tables \\ref{tbl:model_no} and \\ref{tbl:model_sg} for model parameters).  In what follows, we summarize the main results of the present work and discuss them in comparison with observations.  \\begin{figure} \\hspace{0.2cm}\\includegraphics[angle=0, width=8.5cm]{fig16.jpg} \\caption{ Relationships (a) between $\\Qb$ and $\\R=a/b$ and (b) between $\\RQ$ and $\\R$ from our galaxy models (various lines) with a Ferrers prolate bar (Eqs.\\ [\\ref{eq:Qbnum}] and [\\ref{eq:Rmaxnum}]) in comparison with observational results (star symbols) of \\citet{com10}. The observed bars are best represented by $\\fbar=0.25$--$0.5$ and $n\\leq1$. \\label{fig:comp_bar}} \\end{figure}  \\emph{1. Bar Strength Parameter} -- We measure the bar strength using the dimensionless parameter $\\Qb$ (Eq.\\ [\\ref{eq:Qb}]) as is usually done in observational studies. For our galaxy models with a Ferrers prolate bar with index $n$, we calculate $\\Qb$ as well as the radius $\\RQ$ where the maximum bar torque occurs, and provide the fitting formulae (Eqs.\\ [\\ref{eq:Qbnum}] and [\\ref{eq:Rmaxnum}]) as functions of $\\fbar$ and $\\R$.  While $\\Qb$ is linearly proportional to $\\R$ and almost linearly to $\\fbar$, $\\RQ$ is a weakly decreasing function of $\\R$ and independent of $\\fbar$. Both $\\Qb$ and $\\RQ$ become smaller for a more centrally-concentrated bar.  Having found the dependence of $\\Qb$ and $\\RQ$ on the other bar parameters in our galaxy models, it is interesting to apply our results to observed barred galaxies.  Recently, \\citet{com10} measured $\\Qb$, $\\RQ$, $\\R$ (or, equivalently, the bar ellipticity), and the bar semimajor axis $a$ for a sample of nearby galaxies that contain nuclear rings, which is by far the most complete sample. Figure \\ref{fig:comp_bar}a plots as star symbols the empirical relation between $\\Qb$ and $\\R$ from \\citet{com10} for galaxies with $1.5\\leq\\R\\leq3.5$. Also plotted are equation (\\ref{eq:Qbnum}) for various values of $n$ and $\\fbar$.   Note that the trend of $\\Qb$ becoming larger for larger $\\R$ in the observational estimates is entirely consistent with the results of our galaxy models.  When approximating the observed bars using Ferrers prolate spheroids, the observational results for $\\Qb$ \\emph{vs.} $\\R$ can be best described by the bar mass fraction of $\\fbar=0.3$--$0.5$ for $n=1$ inhomogeneous bars and $\\fbar=0.25$--$0.35$ for $n=0$ homogeneous bars. The observed relation between $\\RQ/a$ and $\\R$ shown in Figure \\ref{fig:comp_bar}b appears better explained by the $n=0$ homogeneous bars than inhomogeneous bars.\\footnote{Galaxies with $\\RQ/a>1$ have non-axisymmetric torques dominated by outer spiral arms rather than by bars \\citep{com10}.} We note that many uncertainties surround the observational determinations of $\\Qb$ and $\\RQ$ as they rely sensitively on the bulge subtraction, assumptions on the disk scale height and orientation angle, etc., which are quite uncertain (e.g., \\citealt{lau04,lau06,but06}). Given that observational errors involved in the $\\Qb$ (also likely in $\\RQ$) determinations are typically $\\sim20\\%$ (e.g., \\citealt{com09}), the comparison shown in Figure \\ref{fig:comp_bar} suggests that bars in real galaxies are most likely to have a mass fraction $\\fbar=0.25$--$0.5$ of the spheroidal component, and unlikely to be more centrally concentrated than the $n=1$ case.   \\begin{figure} \\hspace{0.2cm}\\includegraphics[angle=0, width=8.5cm]{fig17.jpg} \\caption{Comparison of our numerical results (squares with errorbars) with the observational measurements for the relations (a) between the dust-lane curvature $\\Dalp$ and $\\Qb$ and (b) between the ring radius $\\Rring$ and $\\Qb$.  Open and filled squares are for non-self-gravitating and self-gravitating models, respectively. In (a), star symbols are adopted from \\citet{com09}, while star symbols and filled triangles in (b) are from \\citet{com10} and \\citet{maz11}, respectively. \\label{fig:comp_ring}} \\end{figure}  \\emph{2. Dust Lanes} -- The imposed non-axisymmetric bar potential readily induces dust-lane shocks across which gas in rotation about the galaxy center loses angular momentum significantly and falls radially inward. Dust lanes in a quasi-steady state approximately follow one of $\\xone$-orbits aligned parallel to the bar major axis. The curvature $\\Dalp$ of dust lanes in our models depends primarily on $\\Qb$ in such a way that they tend to be more curved under a weaker bar. This results from the facts that dust lanes are closer to the bar major axis as $\\Qb$ increases and that inner $\\xone$-orbits are more eccentric than outer ones (see Fig.\\ \\ref{fig:orb_comp}). It also depends, albeit less sensitively, on $\\R$ in that a more elongated bar has more straight dust lanes (Eq.\\ [\\ref{eq:dalpfit}]), since $\\xone$-orbits are rounder with smaller $\\R$ when $\\Qb$ is fixed. Dust lanes are not much affected by self-gravity since they have a Toomre stability parameter greater than unity and are characterized by strong velocity shear.  As mentioned earlier, that dust lanes are more straight under a stronger bar potential was first theoretically predicted by \\citet{ath92b} and later confirmed empirically by \\citet{kna02} and \\citet{com09}. In particular, \\citet{com09} measured the dust-lane curvatures in a sample of 55 barred galaxies that contain clear dust lanes in the SDSS DR7 or NED images, and studied a relationship between $\\Dalp$ and $\\Qb$. Figure \\ref{fig:comp_ring}a reproduces their observational results as star symbols in comparison with our numerical results shown as open and filled squares with errorbars for non-self-gravitating and self-gravitating models, respectively. Note that $\\Dalp$ decreases with increasing $\\Qb$ in both numerical and observational results, although $\\Dalp$ in the simulations corresponds roughly to the lower envelope of the observational results. Note also that our numerical results are unable to reproduce the scatter seen in the observations. These quantitative differences of $\\Dalp$ between our numerical work and the work of \\citet{com09} are likely due to the differences in the spatial ranges of dust lanes where $\\Dalp$ is measured. \\citet{com09} considered a constant-curvature range that varies from galaxy to galaxy, while we fix the position angles of both ends of the range. Also, the fact that the parameter space covered by our models is very limited (i.e., fixed sound speed, $n=1$ Ferrers bar potential, no magnetic field, etc.) may also be partly responsible for the differences between our numerical results and the observations.  \\citet{com09} further noted that there is a large spread of $\\Dalp$ for given $\\Qb$. To find the origin of the spread, they ran a large number of numerical simulations with differing bar and bulge parameters, and found that the spread in $\\Dalp$ can be reduced if the aspect ratio $\\R$ is considered together with $\\Qb$ when fitting the curvatures.   We similarly found that a linear combination of $\\Qb$ and $\\R$ provides a better fit than $\\Qb$ alone, although $\\Dalp$ decreases with increasing $\\R$ in our fit (Eq.\\ [\\ref{eq:dalpfit}]), while $\\Dalp$ is an increasing function of $\\R$ in their fit (Eq.\\ [3] of \\citealt{com09}). This discrepancy is again thought of as arising from the differences in the ranges of dust lanes where $\\Dalp$ is measured and from the limited range of the parameter space in our models, as mentioned above.  \\emph{3. Nuclear Rings: Size} -- The shocked gas moving in toward the central regions along the dust lanes has so large a speed that the bar torque cannot stop its motion across the ILR. The inflowing gas keeps moving in and eventually forms a nuclear ring at the location where the centrifugal force balances the external gravitational force. The mean radius of a ring in our models is generally smaller than the ILR location, and decreases systematically with increasing $\\Qb$ (Eqs.\\ [\\ref{eq:Rring_a}] and [\\ref{eq:Rring_ILR}]). By making the total gravitatioal potential deeper, self-gravity makes the ring larger in size by $\\sim5-20\\%$, with larger values corresponding to lager $\\Qb$. Combined with the results of Paper I that showed that the ring position is insensitive to the mass of a central BH (and thus to the number and locations of ILRs), this clearly evidences that the ring position is not determined by the resonant interactions of gas with the underlying gravitational potential but rather by the amount of angular momentum loss at the dust-lane shocks.  \\citet{com10} also presented the sizes of nuclear rings in their sample of barred spiral galaxies.  Figure \\ref{fig:comp_ring}b compares their results (star symbols) with our simulation outcomes (squares with errorbars) on the $\\Rring/a$--$\\Qb$ plane. Both observational and numerical results show that $\\Rring/a$ becomes smaller with increasing $\\Qb$, indicating that stronger bars can possess smaller rings. For $\\Qb\\simlt0.15$, the agreement between observational and numerical results is quite good. For $\\Qb\\simgt0.15$, on the other hand, $\\Rring/a$ in our models corresponds roughly to the upper envelope of the observational results for given $\\Qb$. This is presumably because our numerical models are unmagnetized: inclusion of magnetic fields efficiently removes angular momentum further at dust-lane shocks, which makes the ring size smaller by a factor of $\\sim2$ when magnetic fields have an equipartition strength with the thermal energy \\citep{ks12}. A larger effective sound speed can additionally make the rings smaller (Paper I).  More recently, \\citet{maz11} measured the ring radii for a sample of 13 barred/unbarred galaxies that contain star-forming nuclear rings: their results for 8 galaxies whose bar sizes are given in \\citet{com10} are plotted in Figure \\ref{fig:comp_ring}b as triangles. \\citet{maz11} argued that the ring size is well correlated with the compactness $\\C\\equiv v_0^2/r_t$, where $r_t$ is the turnover radius of the rotation curve that has the velocity $v_0$ at the flat part, such that more compact (with smaller turnover radius) galaxies have a smaller ring. In our galaxy models, however, $r_t\\sim1\\kpc$ and $\\C\\sim 4\\times10^4 \\;(\\rm km\\;s^{-1})^2 \\kpc^{-1}$, insensitive to $\\Qb$ and $\\R$, while the variation of the ring size with $\\Qb$ is about by a factor of $\\sim2$--$3$. Since the sample galaxies in \\citet{maz11} also exhibit a positive correlation between $\\C$ and $\\Qb$, the negative correlation between $\\C$ and $\\Rring$ in their results may simply be a reflection of a more intrinsic negative correlation between $\\Rring$ and $\\Qb$.  Of course, our current models with a limited range of the parameters cannot address the effects of $\\C$ on the ring size.   \\emph{4. Nuclear Rings: Shape} -- Not all the rings have a conventional shape similar to $\\xtwo$-orbits in our models unless self-gravity is included. In non-self-gravitating models with $\\fbar = 0.6$ and $\\R\\geq 2.0$,  the inflowing gas cannot settle on $\\xtwo$-orbits since they have too small kinetic energy.  The ring material instead takes on an inclined orbit in between $\\xone$- and $\\xtwo$-families. An inclined ring in these models becomes smaller and more eccentric as low angular-momentum gas is added from outside.  It loses much of its mass through the inner boundary when its short axis moves close to the center.  Due to the bar torque, the inclined ring precesses slowly to align its long axis parallel to the bar major axis, eventually forming an $\\xone$-type ring. Strong self-gravity of an inclined ring provides additional non-axisymmetric torque that prevents the precession of an inclined ring, leading instead to an $\\xtwo$-type ring in self-gravitating models.  When $\\fbar \\leq 0.3$ or $\\R\\leq 1.5$, on the other hand, the $\\xtwo$-family of closed orbits near the center have sufficiently large kinetic energy that the inflowing gas along the dust lanes can transit easily to one of them in both non-self-gravitating and self-gravitating models. These rings are eccentric with an ellipticity of $\\ering\\sim0.2$--$0.3$ insensitive to $\\Qb$, which is within the range of the observed ring ellipticities, $\\sim0$--$0.4$, reported by \\citet{com10} and \\citet{maz11}. Again, magnetic fields are expected to circularize nuclear rings \\citep{ks12}.   \\emph{5. Nuclear Spirals} -- Well-defined twin-armed nuclear spirals grow only in models in which the nuclear ring is of the $\\xtwo$-type and sufficiently large in size: they would otherwise be destroyed by the ring material on highly eccentric orbits.  Even in models with an $\\xtwo$-type ring, nuclear spirals are absent if the ring is so small to limit their spatial extent. All nuclear spirals in our models are trailing and logarithmic in shape. While the shapes of dust lanes and nuclear rings do not change much with time after the potential is fully turned on, nuclear spirals are not stationary over the course of the entire evolution. They initially start out as being tightly wound and weak, and then gradually unwind and become stronger until turning into shocks. This unwinding and growth of nuclear spirals appears to be a generic property of nonlinear waves that become more nonlinear as they propagate inward \\citep{lee99}. The unwinding rate is lower in self-gravitating models than in non-self-gravitating models. Since nuclear spirals grow and unwind faster as $\\Qb$ increases, the probability of having more tightly-wound and weaker spirals is larger for galaxies with a weaker bar torque. This is consistent with the observational results of \\citet{pee06} who found that tightly wound spirals are found primarily in weakly barred galaxies, while loosely wound spirals are more common in strongly barred galaxies (see also \\citealt{mar03a,mar03b}). \\citet{pee06} also found that grand-design nuclear spirals in strongly barred galaxies does not extend all the way into the nucleus, which is consistent with our numerical results that show that nuclear spirals cease to exist by turning into shocks, and this happens earlier in higher-$\\Qb$ models.   \\emph{6. Mass Inflow Rates} -- In our models, the mass inflow rate to through the inner boundary is found to be $\\Mdot\\sim10^{-3}-10^{-2}\\Aunit$ for models with $\\Qb\\simlt0.2$ regardless of the presence of self-gravity. Without self-gravity, models with $\\Qb\\simgt0.2$ but with an $\\xtwo$-type ring still have $\\Mdot < 10^{-2}\\Aunit$, since most of the inflowing gas is trapped in the ring.  If the inflowing gas moves all the way to a central BH, these values of $\\Mdot$ correspond to the Eddington ratio $\\lambda\\equiv L_{\\rm bol}/L_{\\rm Edd} = 4.5\\times10^{-2} (\\Mdot/10^{-2}\\Aunit)(\\MBH/10^7\\Msun)^{-1} \\sim 10^{-3}$--$10^{-2}$ (Paper I), potentially explaining low-luminosity Seyfert 1 AGNs. (e.g., \\citealt{ho08}). Here, $L_{\\rm bol}$ and $L_{\\rm Edd}$ denote the bolometric and Eddington luminosities of an AGN, respectively, and 10\\% of the mass-to-energy conversion efficiency of the accreted material is assumed. Some of our numerical models exhibit unrealistically large mass inflow rates. Non-self-gravitating models in which the central regions are dominated by the gas on $\\xone$-orbits are found to have $\\Mdot \\simgt 1\\Aunit$ and sometimes as large as $\\sim10\\Aunit$ when the ring gas on eccentric orbits is accreted directly to the inner boundary. In self-gravitating models with $\\Qb\\simgt0.2$, on the other hand, rings are unstable to form high-density clumps with mass $\\sim10^6-10^7\\Msun$.  These clumps sometimes plunge into the central hole, causing the mass inflow rate to fluctuate with large amplitudes.  We finally remark some caveats associated with $\\Mdot$ obtained in our simulations and in interpreting it as a mass accretion rate to a central BH. First, $\\Mdot$ from self-gravitating models are definitely more realistic that that from non-self0gravitating models. Still, Models M60R30G and M60R35G can not be applied to real galaxies since they do not posses well-defined nuclear rings: the bar in these models is perhaps too massive or too elongated, or the Ferrers prolate bar is not a good representation of realistic bars. Second, as mentioned earlier, $\\Mdot$ measured is the rate of gas mass that goes in through the inner boundary.  This is likely to be an upper limit to the real accretion rate to the BH since some of the inflowing mass changes its orbit before reaching the BH and can possibly come out of the inner boundary. Third, high-density clumps produced by gravitational instability of nuclear rings would undergo star formation, reducing gas content in the rings. Ensuing feedback would destroy them, so that the accretion of dense clumps in real situations would much less frequent than in our simulations. Fourth, a circumnuclear disk with starburst activities can make the accretion rate to the BH much smaller than $\\Mdot$ (e.g, \\citealt{dav07,wat08,kaw08}). Fifth, while magnetic fields are known to enhance $\\Mdot$ considerably \\citep{ks12}, they would suppress or reduce gravitational instability of nuclear rings, which tends to reduce $\\Mdot$. In order to properly evaluate the mass accretion rates, therefore, it is required to run more realistic models of barred galaxies including star formation, feedback, magnetic fields, and other physical processes that affect gaseous features and gas inflows in the nuclear regions."
    },
    "1208/1208.1010_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.1in}} The nonrelativistic annihilation of Majorana dark matter in the Sun to a pair of light fermions is chirality-suppressed. Annihilation to 3-body final states $\\ell^+f^-V$, where $V=W,Z,\\gamma$, and $\\ell$ and $f$ are light fermions (that may be the same), becomes dominant since bremsstrahlung relaxes the chirality suppression. We evaluate the neutrino spectra at the source, including spin and helicity dependent effects, and assess the detectability of each significant bremsstrahlung channel at IceCube/DeepCore. We also show how to combine the sensitivities  to the dark matter-nucleon scattering cross section in individual channels, since typically several channels contribute in models. ",
        "introduction": "Dark Matter (DM) particles $X$ can become captured and trapped at the center of the Sun and the Earth. As the DM density grows over time, the accumulated DM can annihilate and produce a neutrino flux that is observable at a detector on Earth~\\cite{bib:early_solar_dm_refs}. The annihilation channels and the final state decay products are determined by details of physics beyond the Standard Mode (SM).  While it is often assumed that 2-body annihilations dominate, 3-body final states can be the leading contribution in models in which the DM candidate is a Majorana fermion. Such candidates arise, for example in supersymmetric models where the lightest supersymmetric particle is a neutralino. In such models, the cross section for dark matter annihilation to light fermions is severely suppressed~\\cite{bib:other_3body_papers}.  The wavefunction of an initial state consisting of a pair of identical fermions % must be totally antisymmetric, implying either $L=S=0$ or $L=S=1$, where $L$ and $S$ are the orbital angular momentum and spin of the pair, respectively. For the latter case, the annihilation matrix element ${\\cal M}(XX \\rightarrow f\\bar f)$ is necessarily $p$-wave suppressed, and is thus proportional to $v \\ll 1$, where $v$ is the relative velocity of the DM particles. The $s$-wave initial state is CP-odd and has zero total angular momentum; if CP-violating effects are negligible, this state must annihilate to an $L=0$, $S=0$ final state. As the final state fermions $f \\bar f$ emerge back-to-back, they must possess the same helicity. Since particles and antiparticles of the same handedness arise from different Weyl spinors, $s$-wave $XX \\rightarrow f \\bar f$ annihilation requires that the initial state couple to both the $f_L$ and $f_R$ spinors, {\\it i.e.}, a mixed coupling to both L/R chirality. (For further elaboration of these issues see the Appendix of Ref.~\\cite{append}.)  While fermion mass readily provides L-R mixing, it leads to a matrix element suppressed by $m_f / m_X$. If the mass term is the only source of helicity mixing, the 2-body $XX \\rightarrow f\\bar f$ annihilation cross section is heavily suppressed. However, a 3-body final state containing an additional vector boson (VB) can be CP-even with vanishing total angular momentum, even if both fermions arise from the same Weyl spinor. As a result, the 3-body annihilation cross section is not suppressed by $m_f^2 / m_X^2$, and becomes significant despite the additional coupling factor ($\\sim \\alpha$).  Radiative electroweak corrections to DM annihilation were recently considered for the gamma ray~\\cite{Barger:2009xe,Bell:2011if,Barger:2011jg}, positron\\ct{Bergstrom:2008gr,Kachelriess:2009zy} and antiproton~\\cite{Kachelriess:2009zy,Garny:2011cj,Garny:2011ii} spectrum of the annihilations. Recently, solar DM signals from electroweak bremsstrahlung were investigated in Ref.~\\cite{Bell:2012dk}.  In comparison, we consider each significant annihilation channel separately; the corresponding event rates can be summed using annihilation branching ratios which depend on the details of a specific model. We also consider DM annihilation to left-handed and right-handed fermions separately. This is important, as the shape of the neutrino injection spectrum depends significantly on the helicity of the fermions (and in particular on their decay spectra). We also numerically propagate the neutrinos through the Sun and vacuum, with oscillations, scattering and $\\tau$-regeneration fully simulated.  In Section~\\ref{sect:cross_section}, we describe the model we adopt and compute the doubly differential 3-body annihilation cross sections. The injection spectra are presented in Section~\\ref{sect:injection}, and a description of neutrino detection in Section~\\ref{sect:result}. In Section~\\ref{sens}, we investigate the discovery potential of the annihilation channels individually and in combination at the IceCube/DeepCore (IC/DC) detector. We conclude in Section~\\ref{conc}. ",
        "conclusions": "\\label{conc}  If dark matter is a Majorana fermion, its annihilation in the Sun to Standard Model fermions is both chirality and velocity-suppressed. Then, the annihilation may primarily be through 3-body processes, with the emission of a gauge boson.  The neutrino spectra from such channels can differ dramatically from the spectra from $2\\rightarrow 2$ processes.  Dark matter couplings to left-handed leptons necessarily open-up 3-body annihilation channels in which neutrinos are produced directly.  Moreover, the branching fractions to these channels are usually large. The neutrino injection spectra are typically hard, providing for interesting detection possibilities at neutrino detectors.  We considered a model in which $SU(2)$-singlet DM couples to SM leptons (either left-handed or right-handed) via exchange of a new scalar $\\eta_{L,R}$.  For this model, we calculated the 3-body differential annihilation cross sections and the neutrino injection spectra with a full treatment of the helicity correlations of the gauge boson and  $\\tau$ lepton decays. We determined the muon event rates at IceCube/DeepCore arising from each annihilation channel, accounting for neutrino propagation effects, including oscillation, scattering and regeneration.    We calculated $3\\sigma$ sensitivities of IC/DC to the DM-nucleon scattering cross section for several 3-body channels. The different channels are of varying utility in constraining dark matter models; channels with primary neutrinos lead to the best sensitivity.  We then showed how to combine the sensitivities in individual channels, to obtain the sensitivity for a combination of channels as may arise in models.   \\vskip .2in {\\bf Acknowledgments.} D.M. thanks the University of Hawaii for its hospitality during the initial stages of this work. J.K. and D.M. thank the Center for Theoretical Underground Physics and Related Areas (CETUP* 2012) in South Dakota for its support and hospitality during the completion of this work. This research was supported in part by DOE grants~DE-FG02-04ER41291, DE-FG02-04ER41308 and DE-FG02-96ER40969, and by NSF grant PHY-0544278.   \\appendix \\label{appendix}"
    },
    "1208/1208.3824_arXiv.txt": {
        "abstract": "{} {Scintillation noise is a major limitation of ground base photometric precision.} {An extensive dataset of stellar scintillation collected at 11 astronomical sites world-wide with MASS instruments was used to  estimate the scintillation noise of large telescopes in the case of fast photometry and traditional long-exposure regime.} {Statistical distributions of the corresponding parameters are given. The scintillation noise is mostly determined by turbulence and wind in the upper atmosphere and comparable at all sites, with slightly smaller values at Mauna Kea and largest noise at Tolonchar in Chile. We show that the classical Young's formula under-estimates the scintillation noise.The temporal variations of the scintillation noise are also similar at all sites, showing short-term variability at time scales of 1 -- 2 hours and slower variations, including marked seasonal trends (stronger scintillation and less clear sky during local winter). Some correlation was found between nearby observatories.} ",
        "introduction": "One of the main characteristics of astronomical objects is their brightness in different spectral bands. The standard precision of ground-based photometry is adequate in most cases, but a number of astronomical problems require an even greater precision \\citep{Heasley1996,Everett2001}. One of the fundamental factors limiting the precision of ground-based photometry is the stellar scintillation occurring in the atmosphere as a result of its turbulent nature.  The fluctuations of the refractive index cause phase distortion in a plane light wave passing through the atmosphere to an entrance aperture of the telescope. As the wave propagates, the phase distortions lead to a redistribution of the amplitudes between different parts of the wavefront. Averaging within the aperture reduces the fluctuations caused by this mechanism, but does not eliminate them completely. In photometric practice, this effect is considered as an additional source of error. The scintillation noise expressed in stellar magnitudes does not depend on the object's brightness, therefore it cannot be reduced by observing brighter stars.  Scintillation noise in high-precision and fast photometry has been studied for quite a long time \\citep{Young1967,Young1969,Dravins1997a}. The large number of methods proposed for reducing scintillation noise \\citep{Heasley1996,Dravins1998,Gilliland1993,Osborn2010} shows that there is no perfect solution and that further work is needed. Moreover, this same problem is often presented as an argument to do precise photometry from space.  Stellar scintillation is of interest not only to photometry. This phenomenon is a powerful tool for remote sensing of optical turbulence (OT) in the atmosphere. The mechanisms of scintillation generation and its characteristics are well studied theoretically and experimentally because they are closely related to the most important characteristic of the OT above astronomical sites, the seeing.  Unlike the seeing, scintillation received relatively little attention in the astro-climatic work; the error budget of high-precision photometry is still evaluated using the data from \\citep{Young1967} or other disparate estimates. Only recently such studies have been conducted in the general context of the characterization of optical turbulence above different astronomical observatories and prospective sites \\citep{Kenyon2006,2011AstL}. These data are useful for comparing sites in the context of photometry. Moreover, the scintillation noise of long-exposure photometric measurements depends not only on the intensity of high-altitude OT, but on the wind speed at altitudes above the tropopause \\citep{2011AstL}, which is important for the global dynamics of the atmosphere.  This paper presents scintillation noise measurements with the MASS instrument at observatories situated in different geographical areas. Sect.~\\ref{sec:theory} recalls the theoretical description of the scintillation in three basic measurement regimes. In Sect.~\\ref{sec:decomposition} the method used to estimate relevant parameters from the raw data is described. The next Section describes the sites studied here, the original data, and the procedure for calculating the scintillation noise. Results and comparative analysis are presented in Sect.~\\ref{sec:results_s3}, Sect.~\\ref{sec:temporal} describes the temporal variability of scintillation noise, and the final section is a discussion and comparison with other available data. ",
        "conclusions": "The resulting characteristics of the scintillation noise in the regime of short exposures are shown in Table~\\ref{tab:comp2}, for long exposures -- in Table~\\ref{tab:compar}. Instead of the squared parameters, those tables list $S_2$ and $S_3$. The $S_3$ was calculated as the average between the values obtained under the assumptions of longitudinal and transverse wind. For the Armazones site, the average values computed with a weight proportional to the number of measurements for campaigns Armazones~A and Armazones~B are presented.  Recall that the scintillation noise at the zenith $\\sigma_\\mathrm{S}$ and $\\sigma_\\mathrm{L}$ for a telescope with diameter $D$ is calculated by the formulae \\begin{equation} \\sigma_\\mathrm{S} = S_2\\,D^{-7/6} \\end{equation} for short exposures, and \\begin{equation} \\sigma_\\mathrm{L} = S_3\\,D^{-2/3}\\tau^{-1/2} \\end{equation} for a long exposure $\\tau$.  These expressions follow directly from the Eqs.~(\\ref{eq:short_s2}) and (\\ref{eq:long_s3}). To convert the values of the Table~\\ref{tab:comp2} and \\ref{tab:compar} into stellar magnitudes, they should be multiplied by the constant 1.086, i.e., the amplitude of the scintillation noise in magnitudes is $\\sigma_\\mathrm{mag} = 1.086\\,\\sigma$.  As we pointed in Sect.~\\ref{sec:theory}, for the SE regime the central obscuration effect should be considered. A good approximation is presented in \\citet{2012bMNRAS}. For extra-large telescopes ($D \\gtrsim 8$~m), the effect of the turbulence outer scale is substantial in both the SE and LE regimes, and should also be taken into account \\citep{2012bMNRAS}.  The results shown in these tables are in good agreement with the estimates of scintillation parameters $S_2$ and $S_3$ for Tololo and Pach\\'on obtained in \\citet{Kenyon2006} by numerical calculation of the moments $\\mathcal M_2$ and $\\mathcal Y_2$ by Eqs.~(\\ref{eq:s2m2}) and (\\ref{eq:s3y2}) on the basis of the measured OT vertical profiles and the modelled wind profile. Our method gives for the Maidanak observatory a similar estimate of $S_3 = 0.0030$ \\citep{2011AstL}. Unfortunately, other statistically reliable estimates of the scintillation noise are not available to our knowledge.  The results indicate that scintillation noise is mostly defined by the global turbulence at altitudes of 10--15\\,km, at the tropopause and above. The only site in our list where the local effects are noticeable and make a difference with other sites is the Tolonchar. In this case, the proximity to the main ridge of the Andes which is quite high and perpendicular to the global air circulation could create a quasi-stationary vertical vortex with a scale of tens of kilometres.  The stronger scintillation noise (about 10\\% excess) at the SPM observatory is likely a result of a biased seasonal distribution of the observations together with a significant variability from night to night (see Table~\\ref{tab:night3}). On the other hand, the most probable value of $S_3$ at this observatory exceeds the value at Paranal by only 3\\%.  \\begin{table} \\caption{Comparison of the $S_2$ distribution at all sites (the units are $\\mbox{m}^{7/6}$). \\label{tab:comp2}} \\centering \\begin{tabular}{lrrrrrrrr} \\hline\\hline & \\multicolumn{3}{c}{$S_2$ quartiles}\\\\ & 25\\% & 50\\% & 75\\% \\\\ \\hline Armazones & 0.0108 & 0.0132 & 0.0166 \\\\ La Chira & 0.0127 & 0.0155 & 0.0198 \\\\ Mauna Kea & 0.0088 & 0.0109 & 0.0139 \\\\ Pach\\'on & 0.0124 & 0.0155 & 0.0197 \\\\ Paranal & 0.0117 & 0.0144 & 0.0181 \\\\ S.~Pedro Martir & 0.0117 & 0.0147 & 0.0192 \\\\ Shatdzhatmaz & 0.0107 & 0.0132 & 0.0168 \\\\ Tolar & 0.0113 & 0.0136 & 0.0169 \\\\ Tololo & 0.0122 & 0.0150 & 0.0191 \\\\ Tolonchar & 0.0116 & 0.0146 & 0.0191 \\\\ Ventarrones & 0.0120 & 0.0150 & 0.0190 \\\\ \\hline\\hline \\end{tabular} \\end{table}  \\begin{table}[h] \\caption{Comparison of the $S_3$ distribution at all sites (the units are $\\mbox{m}^{2/3}\\,\\mbox{s}^{1/2}$). \\label{tab:compar}} \\centering \\begin{tabular}{lrrrrr} \\hline\\hline & $S_3$ Mode & \\multicolumn{3}{c}{$S_3$ quartiles} \\\\ & & 25\\% & 50\\% & 75\\% \\\\ \\hline Armazones & 0.00300 & 0.00266 & 0.00329 & 0.00410 \\\\ La Chira & 0.00328 & 0.00294 & 0.00357 & 0.00439 \\\\ Mauna Kea & 0.00265 & 0.00238 & 0.00290 & 0.00358 \\\\ Pach\\'on & 0.00298 & 0.00265 & 0.00325 & 0.00402 \\\\ Paranal & 0.00311 & 0.00275 & 0.00336 & 0.00410 \\\\ S.~Pedro Martir & 0.00310 & 0.00280 & 0.00352 & 0.00453 \\\\ Shatdzhatmaz & 0.00280 & 0.00251 & 0.00320 & 0.00409 \\\\ Tolar & 0.00329 & 0.00290 & 0.00346 & 0.00414 \\\\ Tololo & 0.00303 & 0.00265 & 0.00324 & 0.00396 \\\\ Tolonchar & 0.00353 & 0.00306 & 0.00386 & 0.00488 \\\\ Ventarrones & 0.00320 & 0.00272 & 0.00337 & 0.00413 \\\\ \\hline\\hline \\end{tabular} \\end{table}  \\subsection{Comparison of the S3 parameter with the Young's equation} \\label{sec:compar_s3}  Until now, astronomers \\citep[see, e.g.,][]{Everett2001,Mann2011} used the Young's formula \\citep{Young1967} to estimates the contribution of scintillation noise to the accuracy of photometric measurements: \\begin{equation} \\sigma_\\mathrm{L} = 0.0030\\, D^{-2/3} M_\\mathrm{z}^{3/2}\\,e^{-h_\\mathrm{obs}/h_0} \\tau^{-1/2}, \\label{eq:young} \\end{equation} where $D$ is telescope diameter in meters, and $h_\\mathrm{obs}$ is the observatory altitude above sea level. Since the original expression contains a bandwidth rather than exposure, there was a misunderstanding in the translation of one to another. The (\\ref{eq:young}) is taken from \\citep{Gilliland1993}, where it was corrected after the intervention of \\citet{Young1993}. The dependence on the observatory altitude was originally proposed by \\citet{Reiger1963}, who assumed an exponential dependence $C_n^2$ from altitude, but the scale height of $h_0 = 8$~km was established by Young. The numerical coefficient was determined from observations mainly with the 0.9~m telescope.  For comparison with our data, we rewrite the previous formula as $S_3 = 0.0030\\,e^{-h_\\mathrm{obs}/h_0}$. Using this formula and the site altitudes from Table~\\ref{tab:sites}, we obtain $S_3$ estimates from 0.0018 for the highest summit of Mauna Kea to 0.0023 for the Shatdzhatmaz. A comparison with Table~\\ref{tab:compar} shows that the Young's formula underestimates the median amplitude of the scintillation noise by a factor of 1.5 (scintillation power by two times). The values of $S_3$ inferred from the Young's formula are similar to the first quartiles of its actual distributions derived both here and in \\citet{Kenyon2006,2011AstL}.  The dependence of the scintillation on the observatory altitude is ambiguous because of many factors affecting it. Our results do not show this dependence for the most of studied sites located between 2000 and 3000~m a.s.l. because this effect does not exceed the accuracy of the method. Only for the highest site, Mauna Kea, we can possibly relate lower scintillation noise to higher altitude. On the other hand, the Mt.~Ventarrones data show domination of local effects.  \\subsection{Conclusions}  This paper presents the results of the evaluation of scintillation noise in observations on telescopes with large diameter $D \\gg r_\\mathrm{F}$ or $D \\gtrsim 1$~m for the optical and near-infrared in the Earth's atmosphere. These estimates are obtained by an indirect method based on the data of the measurement with the MASS instruments, without involving measurements on large telescopes. However, 1) this method has a reliable theoretical basis, 2) identical instruments were involved and measurements were obtained using the same technique, 3) data for each studied site were obtained over a long time period and have a large statistics.  The scintillation noise at short exposures is quantified by the parameter $S_2$ proposed in \\citet{Kenyon2006}, while the parameter $S_3$ characterizes observations with long exposures. Using these parameters, one can calculate the scintillation noise for a telescope of any reasonable diameter for any long exposure. For example, 8-m telescopes such as Gemini or LSST at Cerro Pach\\'on will have scintillation noise of 50\\,$\\mu$mag for $\\tau = 5$\\,min on a typical night.  The data from seven of the 11 sites studied here should be of particular interest to astronomers because they describe the conditions at the existing observatories Tololo, Pach\\'on, Paranal, S.~Pedro Martir, Mauna Kea, or the observatories soon to come into operation (Armazones, Shatdzhatmaz). Scintillation noise measurements at four other summits in the northern Chile are more interesting from the point of view of general turbulence behaviour in the upper atmosphere.  In addition to the general statistical characteristics of the scintillation noise parameters, their temporal variability was investigated, from the short time scales of minutes to seasonal variations. Characteristics of the short-term (during the night) changes are important for optimizing observational strategy in the high-precision photometry. Usually we can assume that the scintillation noise power is stable enough for a $0.5 - 1$~h, but sometimes it may change significantly on timescales of $10 - 15$~min.  The seasonal variations are significant enough and in some observatories they reach a factor of two in power. They are directly related to the latitude of the observatory. The minimum amplitude is observed for Mauna Kea which is close to the equator and the maximum for the mid-latitude observatory at Mt.~Shatdzhatmaz.  The main conclusion from the comparison of scintillation noise at different observatories is that there are no major differences. For the purpose of choosing the best site for high-precision photometry, all sites are essentially equal. A much larger effect can be achieved by choosing the best season for observations and by using the real-time information about the power of the scintillation in operational planning of the photometric observations.  Apart from the high-precision photometry, the parameters measured in the paper can be interesting for error budget evaluation of high-precision differential astrometry \\citep{Kenyon2006,Cameron2009}."
    },
    "1208/1208.4417_arXiv.txt": {
        "abstract": "A variety of independent observational studies have now reported a significant decline in the fraction of Lyman-break galaxies which exhibit \\Lya emission over the redshift interval $z=6$--$7$.  In combination with the strong damping wing extending redward of \\Lya in the spectrum of the bright $z=7.085$ quasar ULAS 1120$+$0641, this has strengthened suggestions that the hydrogen in the intergalactic medium (IGM) is still substantially neutral at $z\\sim 7$.  Current theoretical models imply \\HI fractions as large as $40$--$90$ per cent may be required to explain these data assuming there is no intrinsic evolution in the \\Lya emitter population.  We propose that such large neutral fractions are not necessary.  Based on a hydrodynamical simulation which reproduces the absorption spectra of high-redshift ($z\\sim 6-7$) quasars, we demonstrate that the opacity of the intervening IGM redward of rest-frame \\Lya can rise rapidly in average regions of the Universe simply because of the increasing incidence of absorption systems which are optically thick to Lyman continuum photons as the tail-end of reionisation is approached.  Our simulations suggest these data do not require a large change in the IGM neutral fraction by several tens of per cent from $z=6$--$7$, but may instead be indicative of the rapid decrease in the typical mean free path for ionising photons expected during the final stages of reionisation. ",
        "introduction": "The opacity observed blueward of rest frame \\Lya in the spectra of distant quasars rises toward higher redshifts (\\citealt{Fan06b,Becker07}), indicating the fraction of neutral hydrogen in the intergalactic medium (IGM) is small but increasing with lookback time.  This observation, coupled with the Thomson optical depth measured from cosmic microwave background data (\\citealt{Komatsu11}), suggests that an extended epoch of hydrogen reionisation was ending by $z \\sim 6-7$.  Reaching further into the epoch of reionisation -- when intergalactic hydrogen is still substantially neutral -- is difficult, but \\Lya selected galaxies may offer one promising route to probing somewhat deeper into this distant era (\\citealt{MiraldaEscudeRees98,Haiman02}).  The increasing \\HI content in the IGM at $z>6$ produces a \\Lya damping wing which can extend redward of a galaxy's \\Lya emission line. The damping wing reduces the visibility of the \\Lya emission (\\citealt{MiraldaEscude98}), and the number of \\Lya emitting galaxies (LAEs) observed in flux-limited surveys will thus decrease as the ambient \\HI fraction rises with increasing redshift (\\citealt{HaimanSpaans99}).  The resonant nature of the \\Lya transmission unfortunately complicates this simple picture.  Source clustering (\\citealt{Furlanetto06}), dust (\\citealt{HansenOh06,Dayal11}), and resonant \\Lya scattering within the sources themselves -- which is sensitive to both the \\HI content and gas velocities (\\citealt{Santos04,Dijkstra10}) -- all impact on the visibility of \\Lya emission. The complex resonant radiative transfer within the galaxies often leads to substantial line redshifts of several hundred $\\rm km\\,s^{-1}$ which can significantly enhance the visibility of LAEs when the surrounding IGM is still substantially neutral (\\citealt{Dijkstra11,Laursen11,Zheng11,Barnes11,Jeeson12}).  Recently, however, a significant decrease in the transmission of \\Lya photons from high redshift galaxies has been found between $z=6$--$7$, (\\citealt{Stark10,Pentericci11,Hayes11,Ono12,Schenker12,CurtisLake12}). This result is based on the rapid decline in the fraction of Lyman break galaxies (LBGs) which exhibit \\Lya emission -- a quantity which is (in principle) less susceptible than the LAE luminosity function (e.g. \\citealt{MalhotraRhoads04,Kashikawa06,Hu10}) to observational selection effects and intrinsic evolution in the LAE population. Based on comparisons to existing numerical simulations and semi-numerical models of reionisation, the reported rapid decrease in the transmission of \\Lya emission at $z\\sim 6 -7$ may be interpreted as a rapid change in the volume-averaged neutral fraction (by several tens of percent) over a rather short redshift interval (\\citealt{Pentericci11,Schenker12,Ono12}).   Such a rapid change in the neutral fraction is, however, at odds with the rather low ionising emissivity suggested by the \\Lya forest data at $z=5$--$6$ (\\citealt{BoltonHaehnelt07b}) which appears to imply a much slower evolution of the neutral fraction.  Theoretical reionisation models matching the \\Lya forest data (e.g. \\citealt{Ciardi12,Kuhlen12,Jensen12}) therefore have considerable difficulty in explaining the rather large neutral fractions ($\\sim 40$--$90$ per cent) that have been inferred from the recent LAE/LBG observations at $z\\sim 7$.  Alternative explanations which -- either individually or in combination -- might allow for a more modest change in the IGM neutral fraction include an increase in the escape fraction of ionising photons, or an increase in the interstellar dust content of the galaxies toward higher redshift (see \\citealt{ForeroRomero12} for further discussion of these points).   In this paper, we shall argue that one other clue to this puzzle is provided by the quasar ULAS J1120$+$0641 at $z=7.085$, recently discovered by the UKIRT Infrared Deep Sky Survey (UKIDSS, \\citealt{Lawrence07}).  The spectrum of ULAS J1120$+$0641 exhibits a strong \\Lya damping wing extending redward of the quasar \\Lya emission line (\\citealt{Mortlock11}) -- exactly what is needed to suppress the generally redshifted (relative to the systemic redshift of the galaxy) emission of LAEs.  Theoretical models predict that the regions surrounding rare, bright quasars are amongst the first to reionise (e.g. \\citealt{Furlanetto04b,Iliev06b,TracCen07,Zahn07,McQuinn07}). If the environment around ULAS J1120$+$0641 is typical of quasar host galaxies at $z>7$, the significantly weaker ionising radiation field expected around even the brightest LAEs implies the impact of the red \\Lya damping wing on LAE visibility should be even stronger.   \\cite{Bolton11} used high resolution simulations of \\Lya absorption to demonstrate that unless a proximate damped \\Lya absorber lies within $\\sim 5\\rm \\,pMpc$ of the quasar, a volume averaged \\HI fraction of at least $\\langle f_{\\rm HI}\\rangle_{\\rm V}=0.1$ is required to reproduce the damping wing observed around ULAS J1120$+$0641.  In this work we show that an \\HI fraction of around 10 per cent could also explain the rapid evolution of the LAEs -- obviating the need for \\HI fractions as large as $40$--$90$ per cent.  The difference in the inferred neutral fraction is due to the inclusion of small-scale ($\\sim 20\\rm\\, pkpc$), optically thick absorptions systems which remain unresolved in large-scale ($\\ga 100 \\rm\\,cMpc$) reionisation simulations.  On including these systems the \\Lya opacity up to a few hundred $\\rm km\\,s^{-1}$ redward of rest frame \\Lya rises more rapidly with \\HI fraction than usually predicted.  As we shall demonstrate, this is because the intervening \\Lya opacity of the IGM at the end of hydrogen reionisation is strongly affected by the increasing size of these systems, which naturally occurs as the metagalactic photo-ionisation rate decreases.  A more modest neutral fraction of $\\sim 10$ per cent is much easier to reconcile with the fact that quasar absorption spectra at $z\\simeq 6$ appear to indicate that the IGM is highly ionised only $180\\rm\\,Myr$ later (\\citealt{Fan06b,Wyithe08}, but see also \\citealt{Mesinger10}).  We begin our analysis in Section 2, where we describe our modelling of the IGM \\Lya opacity and the procedure we adopt for incorporating self-shielded gas into a hydrodynamical simulation.  In Section 3 the properties of the optically thick absorbers in the simulations are discussed in more detail.  We explore the implications of these absorption systems for the visibility of \\Lya emission from galaxies in Section 4, and in Section 5 we conclude.  We assume the cosmological parameters $\\Omega_{\\rm m}=0.26$, $\\Omega_{\\Lambda}=0.74$, $\\Omega_{\\rm b}h^{2}=0.023$, $h=0.72$ and a helium mass fraction of $Y=0.24$ throughout.  Comoving and proper distances are denoted by using the prefixes ``c'' and ``p'', respectively. ",
        "conclusions": "We have used a hydrodynamical simulation to model the \\Lya opacity of the intervening IGM during the final stages of reionisation.  As the photo-ionisation rate drops the opacity redward of rest-frame \\Lya is expected to rise rapidly due to the increasing incidence of optically thick absorption systems.  Our results indicate that the bulk of the \\HI opacity will arise from optically thick systems with column densities $N_{\\rm HI}\\sim 10^{18.5}$--$10^{19.5}\\rm\\,cm^{-2}$.  When including these absorption systems in the simulations, only a moderate rise in the volume averaged neutral fraction is required to significantly reduce the transmission redward of rest-frame Ly$\\alpha$.   This result has an important implication for the interpretation of the recently observed decline in the \\Lya emission from high-redshift galaxies at $z\\sim 6$--$7$ (e.g. \\citealt{Stark10,Pentericci11,Hayes11,Ono12,Schenker12,CurtisLake12}). We find that these observations do not require a large neutral fraction ($\\sim 40$--$90$ per cent) in the intervening IGM at $z=7$ as previously suggested.  Instead, if the rapid decline in the LAE/LBG fraction is further corroborated, it may instead be indicative of the rapid decrease of the mean free path of ionising photons expected at the tail-end of reionisation.  Furthermore, as we find the \\Lya emission from high-redshift galaxies will be suppressed for volume-averaged neutral fractions of only $3$--$10$ per cent, the patchiness of reionisation may produce a more modest impact on the clustering properties of observable high-redshift LAEs than previously predicted.  Our findings may be particularly relevant for future surveys which plan to use the transmission of \\Lya emission and the clustering properties of LAEs to probe deep into the epoch of reionisation. Detecting these LAEs at $z>7$ may well become difficult, even when the IGM is only ten per cent neutral.  It may also mean that spectroscopic confirmation of high-redshift candidate galaxies identified with the drop-out technique at $z>7$ may become problematic, at least until emission lines other than \\Lya can be observed with the James Webb Space Telescope.  Finally, extending theoretical predictions beyond the approximate approach adopted in this work is vital.  This will ultimately require incorporating small-scale absorption systems into simulations which also model the patchy nature of reionisation on large scales.  This remains a considerable but fundamental computational challenge for the current generation of reionisation simulations."
    },
    "1208/1208.6248_arXiv.txt": {
        "abstract": "We have conducted a study of debris disks around F stars in order to explore correlations between rotation, stellar winds, and circumstellar disks. We obtained new  24 \\mum\\ photometry from Spitzer's  Multiband Imaging Photometer for Spitzer (MIPS) camera for a sample of 188 relatively nearby  F dwarfs with various rotation rates and optical colors, and combined it with archival MIPS data for 66 more F stars, as well as Wide-field Infrared Survey Explorer (WISE) data for the entire sample, plus 9 more F stars.   Based on the objects' $K_s-[24]$ and $[3.4]-[22]$ colors, we identify 22 stars in our sample  as having 22 and/or 24 \\mum\\ excesses above our detection limit, 13 of which are new discoveries. Our overall disk detection rate is 22/263, or 8\\%, consistent with previous determinations of disk fractions in the Solar neighborhood.   While fast rotating stars are expected to have strong winds capable of efficiently removing dust, we find no correlation between rotational velocity and infrared excess. Similarly, we find no significant difference in excess detection rate between late-type F stars, which have convective surfaces, and early-type F stars, which have fully radiative envelopes. However, the essentially unknown range of ages in this sample may be washing out any effects relating rotation, winds, and disks. ",
        "introduction": "\\label{sec:intro}  Infrared Astronomy Satellite (IRAS) and Spitzer Space Telescope (Werner \\etal\\ 2004) observations of field A stars and of open clusters have shown that debris disks are common at early ages and decrease in both frequency and luminosity for older stars.   Based on images from the Multiband Imaging Photometer for Spitzer  (MIPS; Rieke \\etal\\ 2004), Su et al.\\ (2006) found about a 30\\% frequency of such disks at 24 $\\mu$m for a sample of $<$ 1 Gyr old A stars, with the upper envelope to the 24 $\\mu$m excess flux decaying roughly as $t_0/t$, with $t_0 \\sim$ 150 Myr.  Studies of a number of open clusters (e.g., Currie et al.\\ 2009 and references therein) agree with the general trends found for the field A dwarfs.  The debris disk frequency for low mass stars is less well-determined, but again observations of field stars (e.g., Carpenter et al.\\ 2009) and of open clusters (e.g., Rebull et al.\\ 2008 and references therein) indicate that the debris disk frequency and luminosity peaks at young ages and declines after that, with few or no debris disks detected at 24 $\\mu$m by Hyades or Praesepe age ($\\sim$ 650 Myr; Cieza et al.\\ 2008, Urban \\etal\\ 2012).  For the relatively luminous A stars, radiation pressure removes dust particles from the disk on timescales $<<$ 1 Myr; the dust in these debris disks must be constantly resupplied from collisional events amongst planetesimals. The lower luminosities of late-type (GKM) dwarfs are less effective in removing the dust via blowout or traditional (radiative) Poynting-Robertson drag. However, as described by Chen et al.\\ (2005) and Plavchan et al.\\ (2005, 2009), winds from rapidly rotating, low mass stars should engender a particulate Poynting-Robertson effect that is more efficient at removing orbiting dust.  F stars inhabit a sweet-spot between the realms of more massive and less massive stars. It is within the F spectral type where outer convective envelopes begin to appear, at spectral type F5. Higher mass stars have convective cores and radiative envelopes, and winds from the photonic Poynting-Robertson effect scour dust from circumstellar disks.  Lower mass stars have strong dynamo activity caused by convective motions and magnetic fields in their outer convective envelopes; strong winds should scour dust from circumstellar disks. As shown by Simon \\& Landsman (1991), Wolff \\& Simon (1997), and others, there is good evidence that the transition to strong, dynamo driven winds occurs for stars later than F5 -- with activity indicators such as \\ion{Ca}{2} HK emission correlated with rotation for later type stars, and not correlated for earlier type stars.  It has long been known that the \\vsini distribution of main sequence stars shows a transition at spectral type F.  Early type (O, B, and A) stars are generally rapid rotators, with average rotational velocities of 200 to 300 km s$^{-1}$.   Main sequence G, K, and M dwarfs (older than a few hundred Myr) are generally very slow rotators, with mean rotational velocities of only a few km s$^{-1}$.  There is thus a sharp drop in the mean rotation rate of main sequence (MS) stars as one goes from F0 to F9.  Interestingly, when this facet of rotation on the MS was first recognized, two possible explanations were advanced -- angular momentum loss via winds (Mestel \\& Roxburgh 1962) and angular momentum redistribution into planetary systems (Huang 1965).  The former explanation was eventually recognized as the primary physical mechanism responsible for the drop in rotational velocities.  However, the strength of those winds -- or at least the amount of angular momentum they carry away -- is a subject of controversy. Some models (e.g., Denissenkov \\etal\\ 2010, Sills \\etal\\ 2000) predict that only the outer convective envelope is spun down (initially) by the wind; because the amount of mass (and the moment of inertia) present in the outer convective envelope of late F stars is very small, in this event the inferred mass loss rate could be small even for rapidly rotating F stars.  Other models, however, predict strong coupling between the convective envelope and the radiative interior, and postulate that the entire star is spun down on short timescales even for late F stars (Bouvier \\etal\\ 1997).  The inferred mass loss rate must be much higher in this event.  Because of the rapid change of wind properties through the F star regime, stars of these masses provide unique laboratories for studies of wind-related effects.   Debris disks could be more common around F stars near the transition between early and late F stars, and could persist for longer periods than for either higher or lower mass stars (and may result in planets with high obliquities; see Winn et al.\\ 2010).  However, for stars with massive disks, the effect of PR drag may be small compared to mutual collisions of dust within a disk such that the grains are ground down into fine enough particles that radiation pressure can remove the dust (e.g., Wyatt 2005).   We have conducted a MIPS 24 $\\mu$m survey of 254 field F dwarfs, half of which rotate quickly ($>$15 km s$^{-1}$), the other half of which rotate slowly ($\\leq$15 km s$^{-1}$). We added to the sample 9 objects detected in WISE but missing from the MIPS archive. Our goal was to determine if (a) F dwarfs have a comparatively high debris disk frequency and (b) we see evidence (via debris disk frequency) for development of winds at the F5 boundary where solar-type winds are expected to develop.  Section~\\ref{sec:analysis} presents the data analysis and reduction, \\S\\ref{sec:disks} identifies infrared excesses, and \\S\\ref{sec:results} discusses the results from this paper. ",
        "conclusions": "\\label{sec:results}  The disk (excess) fractions are listed explicitly in Table~\\ref{excess}.  The blue-rapid sample has the lowest disk fraction at $\\sim$3\\%, and the red-slow has the next lowest disk fraction at $\\sim$6\\%. The blue-slow sample has a disk fraction of $\\sim$8\\%, and the red-rapid sample has the highest disk fraction at $\\sim$14\\%. The errors on these disk fractions can be calculated using the binomial distribution, as per Burgasser \\etal\\ (2003); Table~\\ref{excess} lists these formal errors, and Figure~\\ref{fig:df} represents that information graphically.  Among the entire sample, we find the highest fraction of IR excesses (and inferred disks) in the red-rapid rotators ($\\sim$14\\%).  There are $\\sim$8\\% disks in the sample opposite in  properties, the blue slowly rotating sample.  The lowest disk fraction is the blue rapid sample ($\\sim$3\\%), which is consistent with the blue slow and red slow samples and is perhaps distinct in disk fraction from the red rapid sample ($\\sim$14\\%).   The red-rapid sample may have significantly more disks than the blue-rapid sample, but both blue samples are consistent with each other, and both red samples are (barely) consistent with each other. Both slow samples are consistent with each other, and the two rapid samples may be significantly different from each other.  Taken at face value, these disk fractions are not what we expected, {\\em a priori}, to find.  If winds clear disks, and if early F stars do not have winds, then we would have expected a high disk fraction for both of the blue samples (total of 6/108=6\\%, error range 4-9\\%) compared to the red samples (total of 16/155=10\\%, error range 8-13\\%).   We do not see that; they are even consistent with each other, taking both extremes of the errors.   The division between red and blue stars corresponding to a type of F5 is not particularly subject to debate. The nature of the winds on either side of this boundary almost certainly is different. However, we cannot discern any significant difference in the disk fractions between the red and blue samples.  The division we used between rapid and slow rotators is 15 km s$^{-1}$. As discussed above, this was selected based on where the X-ray flux saturates for G stars, and the sub-samples were constructed to be approximately equally sized, given this cutoff (136 rapidly rotating objects, and 127 slowly rotating objects). Their overall disk fractions are 13/136=10\\% (error 8-13\\%) and 9/127=7\\% (error 5-10\\%). Figure~\\ref{vsini} shows a \\vsini\\ vs.\\ $K_s-[24]$ plot for our whole sample, with our rapid/slow division indicated in the plot. If one instead were to use a division of order 30 km s$^{-1}$, however, this plot appears to show a low excess frequency for rapid rotators and a relatively high excess frequency for slow rotators.  However, the correlation is illusory, in large part because of the distortion of the sub-sample sizes (53 objects with \\vsini\\ $>$ 30 km s$^{-1}$ and 210 objects with \\vsini\\ $\\leq$ 30 km s$^{-1}$).  Nearly 65\\% of the 53 stars with \\vsini $>$ 30 km s$^{-1}$ are early F stars.  Were we to use 30 km s$^{-1}$ instead as the cutoff between rapid and slow, the disk fractions would be 2/34 (6\\%, blue-rapid), 1/19 (5\\%, red-rapid), 4/74 (5\\%,blue-slow), 15/136 (11\\%, red-slow).  The subsamples are so different in size that comparison of them is very difficult, and these disk fractions are identical to what we have already determined for these samples within small number statistics. We will retain our rapid/slow cutoff of 15 km s$^{-1}$ for the remainder of the paper.  The overall disk fraction we obtain here is quite consistent with the disk fraction (at 24 \\mum) obtained elsewhere. A significant complication, however, within our sample and with comparison to other samples, is the ages of the stars, which we now discuss.  Slow rotators with disks are certainly in our sample. A possible explanation for the existence of slow rotators with excesses is long-lived pre-main-sequence circumstellar disks.  A number of theoretical models predict a strong correlation between rotation and primordial disk lifetime (via ``disk locking\"), whereby the angular rotational velocities of PMS stars are magnetically locked to the Keplerian velocity of the inner part of the star's disk as long as accretion continues (e.g., K\\\"{o}nigl 1991, Shu et al.\\ 2000).  Such correlations are observed through mid-IR excesses (taken to arise from primordial circumstellar dust) that are larger around slowly rotating PMS stars (e.g., Rebull et al.\\ 2006, Kundurthy et al.\\ 2006).  These models predict that stars with long-lived accretion disks will arrive on the main sequence as slow rotators and those with short-lived accretion disks will arrive on the main sequence as rapid rotators. If one assumes that stars with long-lived accretion disks are more likely to form well-populated planetesimal belts, and therefore to have detectable debris disks at later ages, then there should be an anti-correlation between \\vsini and 24 \\mum\\ excess for young low mass main sequence stars.  However, whether or not we see this correlation in our data set is certainly intimately tied to the issue of ages in these stars.   While Figure~\\ref{MvBV} demonstrates that there is no large, systematic, and obvious age difference between our rapid and slow rotators, it is almost certainly true that our red, rapidly rotating sample is younger -- on average -- than our red, slowly rotating sample.  This follows simply from the fact that stars in the red sample ($(B-V) >$ 0.43) are expected to have dynamo-driven solar-type winds and should thus spin down as they age. Slow rotators could either be stars (of any age) with long-lived PMS disks that arrived on the ZAMS as slow rotators or older stars that arrived on the MS as rapid rotators but subsequently lost angular momentum due to winds. Rapid rotators must necessarily be relatively young.  Comparison to the rotational velocities of stars in open clusters in this color range -- Pleiades, age $\\sim$100 Myr (Queloz et al.\\ 1998); Hyades, age $\\sim$ 650 Myr (Mermilliod et al.\\ 2009, Figure 4); and M67, age $\\sim$4 Gyr (Melo et al.\\ 2001) -- suggests that the red-rapid sample, and in particular the stars with excesses, probably have ages in the range $<$100 Myr to $\\sim$ 1 Gyr. (In fact, the sole blue rapid rotator with a disk is also consistent with this younger age range, but with only one disk, the significance of this comparison is lessened.) See Figure~\\ref{vsiniandpleiades} for the comparison to the Pleiades.   This suggests that the red rapid rotators are on average somewhat younger than the likely average age for the red slow rotators, but not by a huge amount.  A younger age for the red rapid sample would help explain their higher 24 $\\mu$m excess frequency.  We note that all of the stars with excesses have isochronal age estimates provided in Holmberg et al.\\ (2009).  The ages are all between 1.2 and 4.8 Gyr; the red-rapid sample makes up most of the disks, but is not distinctly younger than the rest of the objects.  These ages seem at variance to the ages inferred from their rotational velocities. For these objects, Holmberg adopts metal poor metallicities ($<$[Fe/H]$> \\approx -$0.2), which may or may not be appropriate.   Spectroscopic metallicity determinations for these stars would help considerably in the interpretation of their ages and hence the origin of their disks.   Nevertheless, the mere existence of significant debris disk excesses around the 11 members of the red-rapid sample is significant. Either the red-rapid sample includes a set of stars that are significantly younger than the rest of our sample (and the high disk frequency simply reflects the fact that younger stars more frequently have debris disks) or the red-rapid sample includes a set of stars of ``normal'' age but unusual dust content. Better ages for these stars would address the first possibility; X-ray fluxes for this set of stars could help constrain whether they have unusually weak winds (for example) by use of the Wood \\etal\\ (2000) formula relating X-ray flux and mass loss rate for low-mass stars.   Unfortunately, no X-ray fluxes are available for the red-rapid rotators with IR excesses, hence we cannot estimate their mass loss rates using the Wood et al.\\ (2000) formula.  X-ray data for these stars, and better age estimates (perhaps obtainable by combining Gaia parallaxes and improved [Fe/H] estimates) could place interesting constraints on the dust production rates for these stars.   Sierchio et al.\\ (2010), Stauffer et al.\\ (2007), and Rebull et al.\\ (2008) have all previously noted a possible correlation between 24 $\\mu$m excess and rotation, with excesses being rare or absent amongst rapid rotators.   The samples of stars used in these studies have included F, G, and K dwarfs, and thus are on average significantly later type than our sample.  Even if the correlation with rotation in those samples is real (the statistical significance was not high), it does not necessarily contradict our result.  The wind mass loss rates from G and K dwarfs are likely much higher than for F dwarfs; the comparison is made difficult because direct measurements are not possible, and inferences from rotational velocities depend on the unknown degree to which the spindown affects just the outer convective envelope or also the radiative core. Moreover, the relative strengths of wind and radiation pressure are important and F stars have higher luminosities than later-type stars (the ratio of mass loss to total luminosity is lower).  If the G and K dwarf mass loss rates are much higher, they could successfully scour their disks of dust and produce the apparent correlation, while the weaker winds of F dwarfs still allow debris dust.  One of our original goals was to determine if F dwarfs have a comparatively high debris disk frequency. By comparison to the FEPS GK stars of similar ages, our F stars have a comparable disk frequency. Another of our original goals was to look for evidence (via debris disk frequency) for development of winds at the F5 boundary where solar-type winds are expected to develop, and we do not see strong evidence of this.  A larger sample of F stars would presumably reveal more disks, and thus enable better restrictions on the observed disk fraction.  Observations in other infrared wavelength bands would also enable us to see if our 24 \\mum\\ excesses extend to longer or shorter wavelengths.  A wider range of observations would allow for more conclusive evidence in understanding what is going on in these stars. It would also be of interest to compare our infrared observations with X-ray observations to see if there is any indication of strong activity, which would help to argue for or against the theory of Chen et al.\\ (2005)."
    },
    "1208/1208.4598_arXiv.txt": {
        "abstract": "Measurements of galaxy cluster kinematics are important in understanding the dynamical state and evolution of clusters of galaxies, as well as constraining cosmological models. While it is well established that clusters exhibit non-spherical geometries, evident in the distribution of galaxies on the sky, azimuthal variations of galaxy kinematics within clusters have yet to be observed. Here we measure the azimuthal dependence of the line-of-sight velocity dispersion profile in a stacked sample of 1743 galaxy clusters from the Sloan Digital Sky Survey (SDSS). The clusters are drawn from the SDSS DR8 redMaPPer catalog.  We find that the line-of-sight velocity dispersion of galaxies lying along the major axis of the central galaxy is larger than those that lie along the minor axis. This is the first observational detection of anisotropic kinematics of galaxies in clusters. We show that the result is consistent with predictions from numerical simulations. Furthermore we find that the degree of projected anisotropy is strongly dependent on the line-of-sight orientation of the galaxy cluster, opening new possibilities for assessing systematics in optical cluster finding. ",
        "introduction": "\\label{sec:introduction} Clusters of galaxies constitute the largest, gravitationally collapsed, structures in the universe. They offer a unique opportunity to study the formation and evolution of structure on cosmological scales. Additionally, the mass distribution of galaxy clusters can be observationally probed both dynamically and via weak lensing, making galaxy clusters ideal laboratories for studying dark matter and modifications to general relativity (e.g. \\cite{2010MNRAS.406.1796R,2012arXiv1205.4679R}).  Observations and simulations clearly show that clusters exhibit triaxial rather than spherical shapes \\citep{1982A&A...107..338B,2005ApJ...627..647B,2007ApJ...664..117G}. \\cite{2004ApJ...606...67H} made the first detection of the flattening of galactic dark matter halos using weak lensing and \\cite{2012arXiv1206.4304V} measured the azimuthal variation of the weak gravitational lensing signal around galaxies, constraining galactic dark matter halo ellipticity. Non-zero ellipticity of dark matter halos has also been confirmed in strong lensing measurements by studying the angular distribution of giant arcs around groups and clusters of galaxies \\citep{2004ApJ...609...50D,2012ApJ...749...38M}.  The velocity fields of galaxies in clusters are known to feature non-trivial radial dependencies, but azimuthal variations are often overlooked. If the shape of the underlying dark matter halo is indeed multiaxial, it is natural to think that it will be reflected in an anisotropic velocity field. Indeed this has been confirmed and studied using numerical simulations \\citep{1997MNRAS.290..411T,2005ApJ...629..781K,2010MNRAS.408.1818W,2012arXiv1203.5708S}. Because cluster mass estimators often assume spherically symmetric velocity fields, characterizing the degree of anisotropy in galaxy clusters is crucial in understanding its impact on cluster mass estimates. It may also provide clues to the ongoing formation of galaxy clusters.  The effect of an anisotropic velocity field should manifest itself in an azimuthal variation of the projected velocity dispersion of member galaxies. To our knowledge no attempts have been made to measure the azimuthal variation of galaxy velocity dispersions in a large sample of galaxy clusters.  The aim of this Letter is to test for azimuthal dependence of the projected velocity dispersion of cluster galaxies in a stacked sample of galaxy clusters. We use the Sloan Digital Sky Survey (SDSS) DR8 \\citep{2011ApJS..193...29A} \\redmapper\\ cluster catalog (described below) to search for an anisotropic velocity field by measuring the velocity dispersion of cluster galaxies along the major axis of the central galaxy (CG), and galaxies along the corresponding orthogonal axis.  The results are compared with expectations from numerical simulations. We adopt a flat $\\Lambda$CDM cosmology with $\\Omega_m=0.3$ and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": "We present the first observational detection of anisotropic kinematics of galaxies in clusters of galaxies. We have measured line-of-sight velocity dispersions of member galaxies along the major and minor axes of the CG, in a large sample of stacked clusters from the SDSS. The projected CG position angle is used to separate galaxies into two section depending on their projected proximity to either axes. Galaxies closer to the projected CG major axis are found to have a preferentially larger velocity dispersion than those that are closer to the minor axis, with an average difference in dispersion of $\\Delta\\sigma_v = 38$ km s$^{-1} \\pm 13$ km s$^{-1}$, corresponding to a relative difference of $\\Delta\\sigma_v / \\langle\\sigma_v\\rangle = 6$\\% $\\pm$ 2\\%. These results are consistent with results from cosmological simulations, with $\\chi_\\text{sim}^2/\\text{dof} = 10.4 / 8 = 1.3$.  The presence of this effect is a signature of a prolate velocity ellipsoid in galaxy clusters. Keeping in mind that velocity ellipsoid is aligned with the halo shape ellipsoid \\citep{2005ApJ...629..781K,1997MNRAS.290..411T,2010MNRAS.408.1818W} and the halos tend to be oriented with the cosmic web \\citep{2006MNRAS.370.1422A,2007MNRAS.375..184B}, the effect is likely a remnant of an anisotropic halo formation (\\citealt{2012JCAP...07..042S}; R. Wojtak et al. 2012, in preparation). This raises an important question as to what extent cluster dynamics is influenced by the infall process and cluster location in the cosmic web. The result can also be interpreted in light of studies indicating preferential alignment between the CG and its host halo \\citep{1982A&A...107..338B, 2008MNRAS.390.1562H, 2010MNRAS.405.2023N}, strengthening the case for anisotropic cluster formation.  We furthermore find a clear correlation between the degree of projected anisotropy and the line-of-sight orientation of cluster halos in the simulations. This result opens exciting new possibilities for assessing the degree of random alignment of galaxy clusters. Specifically it might be applied to a sample of optically selected clusters to evaluate the level of selection bias."
    },
    "1208/1208.5885_arXiv.txt": {
        "abstract": "{Coronal magnetic null points have been implicated as possible locations for localised heating events in 2D models. We investigate this possibility about fully 3D null points.} {We investigate the nature of the fast magnetoacoustic wave about a fully 3D magnetic null point, with a specific interest in its propagation, and we look for evidence of MHD mode coupling and/or conversion to the Alfv\\'en mode.} {A special fieldline and flux-based coordinate system is constructed to permit the introduction of a pure fast magnetoacoustic wave in the vicinity of  proper  and improper 3D null points. We consider the ideal, $\\beta=0$,  MHD equations which are solved using the {\\emph{LARE3D}} numerical code. The constituent modes of the resulting wave are isolated and identified using the special coordinate system. Numerical results are supported by analytical work derived from perturbation theory and a linear implementation of the WKB method.} {\\emph{An initially pure fast wave is found to be permanently decoupled from the Alfv\\'en mode both linearly and nonlinearly  for both proper and improper 3D null points}. The pure fast mode also generates and sustains a nonlinear disturbance aligned along the equilibrium magnetic field. The resulting pure fast magnetoacoustic pulse has transient behaviour  which is found to be governed by the (equilibrium) Alfv\\'en-speed profile, and a refraction effect focuses all the wave energy towards the null point. } {Thus, the main results from previous 2D work do indeed carry over to the fully 3D magnetic null points and so we conclude that {\\emph{3D null points are    locations for preferential heating in the corona by 3D fast magnetoacoustic waves}}.} ",
        "introduction": "\\label{section:1} An abundance of observational data from solar instruments including {\\emph{SOHO}} (e.g. Ofman et al. \\cite{ofman.etal.1997}; DeForest \\& Gurman \\cite{plumes}), {\\emph{TRACE}} (e.g.  Nakariakov {{et al.}} \\cite{Nakariakov1999}; De Moortel et al. \\cite{demoortel.etal.2000}) and {\\emph{Hinode}} (e.g. Ofman \\& Wang \\cite{ofman.wang.2008}) confirms the existence of magnetohydrodynamic (MHD) wave motions in the coronal plasma (see also reviews by De Moortel \\cite{ineke2005};  Nakariakov \\& Verwichte \\cite{Nakariakov2005}; Ruderman \\& Erd\\'elyi \\cite{Misha2009};    Goossens et al. \\cite{goosens2011}; De Moortel \\& Nakariakov \\cite{ineke2012}). MHD wave theory suggests that the corona potentially supports a wide variety of distinct classes of wave motions, namely the Alfv\\'en wave, and fast and slow magnetoacoustic waves. However, the applicability and appropriateness of such classifications has recently provoked intense discussion, for instance reports of having  observed Alfv\\'en waves in the corona (Tomcyzyk et al. \\cite{tomczyk.etal.2007}; De Pontieu et al. \\cite{Bart2007}) are contested by, e.g., Erd\\'elyi \\& Fedun (\\cite{Erdelyi2007}) and Van Doorsselaere et al. (\\cite{Tom2008}).    The \\lq{classic}\\rq{ } terminology of MHD waves originates in the analysis of modes supported by magnetically-unidirectional, homogeneous plasmas of infinite extent which consider either plane wave solutions or utilise the method of characteristics (viz. Riemann decomposition) (examples of such analyses can be found in, e.g. Friedrichs \\& Kranzer \\cite{friedrichs}; Lighthill \\cite{lighthill}; Cowling \\cite{Cowling1976}; Goedbloed \\& Poedts \\cite{goedbloedbook1}). Here, three distinct modes are permitted, the Alfv\\'en mode, and the fast and slow magnetoacoustic, and their behaviour and nature is well understood. For a low-$\\beta$ plasma, it is found that the Alfv\\'en wave is a transverse, purely-magnetic wave, propagating at the Alfv\\'en speed, guided by the magnetic field. The fast magnetoacoustic wave is found to propagate roughly isotropically  at the fast speed ($c_F=\\sqrt{c^{2}_{A}+c^{2}_S}$, $c_F\\approx c_A$ where $\\beta\\ll1$), and can travel along and across magnetic fieldlines. The slow magnetoacoustic wave propagates longitudinally along the magnetic fieldlines, roughly at the sound speed.   However in the solar corona,  effects including gravitational stratification, inhomogeneous density profiles and multiple-source magnetic field geometries call into question whether these three, classical modes are still valid. The departure to inhomogeneity typically introduces a variety of phenomena such as resonant absorption (in the corona see, e.g., Ionson \\cite{Ionson78}; \\cite{Ionson82};  \\cite{Ionson83}; Hollweg \\cite{Hollweg1984}; Ruderman \\& Roberts \\cite{rudermanroberts}) and phase mixing (e.g. Heyvaerts \\& Priest \\cite{HP1983}), all of which blur the distinction between these separate modes (see also, e.g.,  Bogdan et al. \\cite{bogdan}; McDougall \\& Hood  \\cite{Dee2007}; Sousa \\& Cunha \\cite{Sousa2008}; Cally \\& Hansen \\cite{Cally2011}; Hansen \\& Cally \\cite{Hansen2012}).   Typically theorists preserve the concept of the three MHD modes and introduce the concept of conversion and coupling between the constituent modes as waves encounter certain inhomogeneous features. However, as highlighted by Goossens et al. (\\cite{goosens2011}), we must take care to understand that realistic plasmas, strictly speaking, do not support three \\lq{classic}\\rq{ } wave modes but rather that a propagating MHD pulse encounters situations in which it assumes transient properties that qualitatively correspond to one (or more) of the homogeneous modes. Conversely, it should also be possible to locally determine hyperbolic characteristics corresponding to the fast, slow and Alfv\\'en waves for any ideal MHD system via Riemann decomposition (for an overview see, e.g., Goedbloed et al. \\cite{goedbloedbook2}) although, as discussed in Goedbloed \\& Poedts (\\cite{goedbloedbook1}), this is not always analytically possible for fully inhomogeneous cases (in particular for non-unidirectional magnetic fields).  As such, it is unclear whether or not the classic MHD modes exist in the solar atmosphere as discrete entities or if such mode separation requires specific geometries. For example, as discussed by Parker (\\cite{parker91}), the existence of true Alfv\\'en waves as per Alfv\\'en (\\cite{alfven}) is dependent upon special magnetic geometries which contain invariant directions. Despite this uncertainty, a theoretical framework which involves them conceptually is not necessarily obsolete. \tMode interpretation and analysis that relates  MHD wave behaviour in inhomogeneous, realistic plasmas to the comparatively simple, classic waves still has the potential to be a useful framework.    In this paper, we are specifically interested with the behaviour of MHD waves in the vicinity of the 3D {\\emph{magnetic null point}} (e.g. Parnell et al. \\cite{Parnell1996}; Priest \\& Forbes \\cite{magneticreconnection2000}). Null points are topological  features of coronal magnetic fields, predicted by magnetic field extrapolations such as Beveridge et al. (\\cite{BPB}) and Brown \\& Priest (\\cite{BP}). At these magnetic null points, the magnetic induction is zero. Hence, approaching a magnetic null point, the coronal plasma is highly inhomogenous. In the solar atmosphere, null points have been identified as playing key roles in many processes, for example; in  CMEs (in the \\emph{magnetic breakout model}, e.g. Antiochos \\cite{Antio98}; Antiochos et al. \\cite{Antio99}), magnetic reconnection (e.g. Priest \\& Forbes \\cite{magneticreconnection2000}) and oscillatory reconnection (reconnection driven by wave-null interactions, see McLaughlin et al. \\cite{MDHB2009}; McLaughlin et al. \\cite{james2012}) all of which are thought to play a role in coronal heating. As both waves and null points are ubiquitous in the corona (Close et al. \\cite{Close04}, Longcope \\& Parnell \\cite{Longcope09} and R\\'egnier et al. \\cite{Regnier08} give a rough estimates of $1.0-4.0\\times10^{4}$ null points) wave-null interactions are inevitable and thus are arguably a fundamental plasma process in the solar atmosphere. Hence, our present research is specifically concerned with the extension of classic MHD wave theory  about null points,  which in a broader sense contributes to both the theory of MHD wave behaviour in inhomogeneous media and the understanding of fundamental aspects of coronal physics.   MHD wave behaviour in the neighbourhood of magnetic null points have been extensively studied in 2D models. Bulanov \\& Syrovatskii (\\cite{Bulanov1980}) performed the first investigation of MHD behaviour about a 2D null and noted that in a 2D geometry the  motions governing the Alfv\\'en mode and the fast magnetoacoustic modes are decoupled, permitting analysis that considers the modes separately.  The transient features of fast and Alfv\\'en waves in various 2D null point geometries within a $\\beta=0$ plasma were studied extensively in a series of papers by McLaughlin \\& Hood (\\cite{MH2004}; \\cite{MH2005}; \\cite{MH2006a}). Again, these authors found the two modes decoupled and identified key propagation features for each mode. The fast wave propagates isotropically (at the same characteristic speed both across and along magnetic field lines), with behaviour dictated by the Alfv\\'en-speed profile, propagating from regions of high to low Alfv\\'en speed resulting in a refraction effect which focuses the wave energy into the null point. Meanwhile, the Alfv\\'en wave is confined to follow magnetic fieldlines, thus leading the wave energy to accumulate along the separatrices. Due to the resulting current build-up in these regions, these papers concluded that magnetic null points are likely locations for localised heating events in the corona. $\\beta \\neq 0$ and nonlinear behaviour has also been investigated by McLaughlin \\& Hood (\\cite{MH2006b}) and McLaughlin et al. (\\cite{MDHB2009}) respectively. A comprehensive overview of the whole topic is given in the review paper of McLaughlin et al. (\\cite{james2011a}).    Of course, 2D models only give an initial grounding in the physics of realistic null points and for a full understanding we must turn to 3D models  (as singularities, null points are inescapably 3D and any 2D X-point configuration in fact only captures the physics of a {\\emph{null line}} of infinite extent). It is not clear to what extent the characteristics and behaviour of waves about 2D null points transfer to the fully 3D  case, and surprisingly few papers have been written that address MHD wave behaviour about a fully 3D null point. Most papers have focused on the dynamics of current accumulation over time (in an attempt to locate regions where reconnection is most likely to occur) rather than  focus on the transient propagation features of the individual modes.      Galsgaard et al. (\\cite{Galsgaard2003}) consider a proper, $\\beta=0$ null point and introduce a twist wave (what they call an Alfv\\'en wave) which is generated about the spine and eventually accumulates on the {\\emph{fan}} (perhaps analogous to Alfv\\'en waves behaviour around 2D null points). In addition, they observe a small amount of current accumulation at the null itself, which they suggest is due to nonlinear generation of a fast wave. No linear coupling between wave modes is observed, and in their linear analysis they find that the wave equations for the fast and Alfv\\'en modes decouple. However, this is not surprising since their proper 3D null has azimuthal symmetry, and so the system is actually 2.5D, not fully 3D.   Pontin \\& Galsgaard (\\cite{PG2007}), Pontin et al. (\\cite{PBG2007}) and Galsgaard \\& Pontin (\\cite{klaus2011a}; \\cite{klaus2011b}) performed numerical simulations  in which the spine and fan of a proper 3D null point are subjected to rotational and shear perturbations. They found that  rotations of the fan plane lead to current density accumulation about the spine,  and rotations about the spine lead to current sheets  in the fan plane. In addition,  shearing perturbations lead to 3D localised current sheets focused at the null point itself. Again, this is in good agreement with what we may expect for MHD wave behaviour from the 2D studies, i.e.  current accumulation at specific, predictable parts of the magnetic topology.   The first study of MHD wave behaviour in the neighbourhood of a fully 3D null point was investigated by McLaughlin et al. (\\cite{MFH2008}). These authors examine the fast and Alfv\\'en waves about both proper (i.e. 2.5D) and improper (fully 3D) null points. The authors utilise the WKB approximation to determine the transient properties of the modes in a linear, $\\beta=0$ plasma regime. Their findings strongly suggest that the features of MHD waves about 3D nulls are not that different to the 2D results: the fast wave propagates across magnetic fieldlines according to the Alfv\\'en-speed profile and the Alfv\\'en wave is confined to magnetic field lines, leading to the waves accumulating in particular topological regions of the null point. However, their implementation of the WKB method is unable to address the question of whether modes couple in the fully 3D geometry, since their first-order WKB solution implicitly precludes the possibility and constrains the waves to see the magnetic field as locally uniform.      Thus, as it stands the question of what is the true behaviour of MHD waves in the neighbourhood of fully 3D magnetic null points remains unanswered. Specifically, what is the nature of the propagation and evolution of each MHD wave and, critically, what is the efficiency of mode-coupling and/or conversion due to the magnetic geometry and/or due to nonlinear effects?   \\begin{figure*}[t] \\centering \\includegraphics[width=17cm]{fig1.eps} \\caption{Left: The azimuthally symmetric proper 3D null point ($\\epsilon=1$). Right: An improper 3D null point where fieldlines are predominantly aligned with the $x$-axis ($\\epsilon=0.5$).} \\label{fig1} \\end{figure*}  This paper investigates the behaviour of fast magnetoacoustic waves about fully 3D null points via numerical simulations and looks for evidence of coupling to the Alfv\\'en mode in any form (e.g. linear, nonlinear and due to the 3D magnetic geometry). The paper is structured as follows: $\\S \\ref{section:2}$ outlines the specifics of the model and is subdivided into sections on the structure of the null points considered ($\\S \\ref{section:2.1}$), the mathematical model ($\\S \\ref{section:2.2}$), the method for isolating individual wave modes ($\\S \\ref{section:2.3}$), the specifics of the numerical solution ($\\S \\ref{section:2.4}$) and a brief outline of the supporting WKB solution ($\\S \\ref{section:2.5}$). $\\S \\ref{section:4}$ and $\\ref{section:3.1}$ present the results of the simulations for the improper and proper radial null points respectively, and the conclusions are presented in $\\S \\ref{section:5}$. ",
        "conclusions": "\\label{section:5}   We have studied both the transient properties of fast magnetoacoustic waves and the nature of linear and nonlinear mode coupling in the vicinity of two potential 3D null points: the proper ($\\epsilon=1$) null, which can be treated as a 2.5D problem using cylindrical polars, and the fully 3D, improper ($\\epsilon=0.5$) null point. Regardless of the eccentricity, $\\epsilon$, of the null point studied, we find that: \\begin{itemize} \\item {An initially pure, linear fast wave exhibits no linear or nonlinear coupling to the Alfv\\'en mode in the neighbourhood of a 3D null point.} \\item {Due to the absence of coupling, the propagation of the fast wave is entirely dictated by the Alfv\\'en-speed profile which about a 3D null point leads to a refraction effect,  focusing all the wave energy at the null point itself.} \\item {The propagating fast wave generates and sustains an instantaneous and dependent nonlinear field-aligned disturbance.} \\end{itemize}  One of the chief aims of the simulations  was to look for evidence of MHD mode coupling about fully 3D (i.e. improper) null points. Unlike the proper \\textit{symmetric} 2.5D null, it  was not clear from analysis of the MHD equations as to whether wave modes about improper 3D nulls are coupled due to the asymmetric magnetic topology, nor whether these coupling terms are linear or nonlinear, \\emph{i.e. it was not clear what the overall general effect of departing from azimuthal symmetry would have on the wave dynamics.} To address this we have generated a pure fast magnetoacoustic wave in the vicinity of 3D nulls and analysed its subsequent propagation. We find that this wave exhibits no mode coupling to the Alfv\\'en mode at either a linear or nonlinear scale about both the proper ($\\epsilon=1$) and improper null point ($\\epsilon=0.5$). Further experiments, namely $\\epsilon=0.25$ and $\\epsilon=0.75$ (values $\\epsilon>1$ have analogues in the range $0<\\epsilon<1$, and $\\epsilon=0$ recovers the 2D null point), retain these results. Thus, we find that lack of mode conversion is a general feature of all potential null points, \\emph{i.e. magnetic fieldline eccentricity does not facilitate fast to Alfv\\'en mode conversion.}  We additionally find that, in both numerical experiments, the fast-mode pulse generates and sustains a field-aligned disturbance in $|\\mathbf{v}_{\\mathbf{B}_0}|$.   Our analysis shows that this field-aligned disturbance is not self-sustaining: it is an instantaneous {\\emph{daughter}} disturbance resulting from the {\\emph{progenitor}} fast wave. It is evident that the disturbance in $|\\mathbf{v}_{\\mathbf{B}_0}|$  does not act as a mode-conversion mechanism in the cases considered here, since no signal is ever generated in $|\\mathbf{v}_\\mathbf{A}|$,  and the disturbance has no feedback effect upon the main wave.  Due to the absence of mode coupling, the key features of fast wave behaviour in the neighbourhood of the $\\epsilon=0.5$, improper null points are fundamentally the same as the behaviour in 2D, $\\beta=0$ null point studies (e.g. McLaughlin \\& Hood \\cite{MH2004}), despite the inhomogeneous, fully 3D magnetic field. The difference between propagating fast waves about different potential null points is due \\textit{only} to the differing Alfv\\'en-speed profiles. Thus, the propagation of the fast wave is found to be entirely dictated by such profiles, causing the refraction effect which, over time, focuses all of the wave energy at the null point. This is confirmed by our numerical simulations and supports the conclusions drawn from a 3D WKB approach in McLaughlin et al. (\\cite{MFH2008}). Note, in the numerical simulations at large time, the pulse is so close to the null point that the resolution will eventually become inadequate. Nonetheless, in our simulations, over a finite, resolvable  time the wave energy accumulates in a small spatial region around the null point, resulting in an exponential steepening of current-gradients, hence resistivity will eventually become non-negligible, resulting in ohmic heating of the local coronal plasma via resistive dissipation (a conclusion directly carried-over from 2D studies, see McLaughlin et al. \\cite{james2011a}). Thus, \\emph{we conclude that 3D null points are locations for preferential heating by passing fast magnetoacoustic waves.}   Wave motions and null points are ubiquitous in the corona and so these fast wave-null interactions  are likely inevitable. The large body of theory, extended by the results presented here, indicates that this will result in localised heating events. Thus, the key question is \\emph{does this make a signifigant contribution to coronal heating?} To answer this, clear observational evidence for MHD waves around coronal nulls is needed, a calculation of resultant heating, and a survey of the prevalence of such events. Such observations would require the detection of coronal nulls using co-temporal high spatial/temporal imaging and magnetograms to study oscillations in the vicinity of coronal null points. It is possible that the \\emph{Atmospheric Imaging Assembly} and \\emph{Helioseismic and Magnetic Imager} aboard SDO may permit such a study.  We cannot, however, assume the  results presented here hold in the opposite case, i.e. a similar scenario for an initially pure Alfv\\'en wave, follows suit. In fact, results of  Galsgaard et al. (2003) report this may not be the case.     Such an investigation will be the subject of a future paper.  Finally,  we highlight that sufficiently close to the null point, as magnetic induction drops off, there will be a region where the magnitudes of the sound speed and Alfv\\'en speed become comparable: identified by Bogdan et al. (\\cite{bogdan}) as the \\emph{magnetic canopy}, \\textit{viz.} the $\\beta=1$ layer. This was investigated for a 2D null point by McLaughlin \\& Hood (\\cite{MH2006b}) and has the effect of introducing the \\textit{slow magnetoacoustic wave} into the system, which is coupled to the fast wave (but both are still decoupled from the Alfv\\'en wave). Thus the cold, $\\beta=0$ plasma assumption does not completely capture the physics of MHD waves about real null points, which \\emph{must} contain a $\\beta=1$ layer near the null point. Thus, to understand wave dynamics about null points in the corona, it is necessary to extend the model presented here and see how these results are modified when the $\\beta=0$ assumption is removed."
    },
    "1208/1208.2183.txt": {
        "abstract": "{ SN 2004et is one of the nearest and best-observed Type IIP supernovae, with a progenitor detection as well as good photometric and spectroscopic observational coverage well into the nebular phase. Based on nucleosynthesis from stellar evolution/explosion models we apply spectral modeling to analyze its $140-700$ day evolution from ultraviolet to mid-infrared. We find a $M_{\\rm ZAMS}= 15$ \\msun~progenitor star (with an oxygen mass of 0.8 \\msun) to satisfactorily reproduce [\\ion{O}{i}] \\wll6300, 6364 and other emission lines of carbon, sodium, magnesium, and silicon, while 12 \\msun~and 19 \\msun~models under- and overproduce most of these lines, respectively. This result is in fair agreement with the mass derived from the progenitor detection, but in disagreement with hydrodynamical modeling of the early-time light curve. %We calculate synthetic light curves in all bands from U to K , and find agreement with a \\iso{56}Ni mass of $\\sim$0.056 \\msun. From modeling of the mid-infrared iron-group emission lines, we determine the density of the ``Ni-bubble'' to $\\rho(t) \\simeq 7\\e{-14}\\times(\\mbox{t}/100~\\mbox{d})^{-3}$ g cm$^{-3}$, corresponding to a filling factor of $f = 0.15$ in the metal core region ($V = 1800$ \\kms). We also confirm that silicate dust, CO, and SiO emission are all present in the spectra.} % context heading (optional) ",
        "introduction": "%\\textbf{Overview Type IIPs}\\\\ Stars with Zero Age Main Sequence (ZAMS) mass greater than about 8 \\msun~end their lives as core-collapse supernovae (SNe). Over half of these events (per unit volume) are classified as Type IIP \\citep{Li2011}, showing hydrogen lines, as well as a $3-4$ month plateau in the light curve, implying the presence of a massive hydrogen envelope. %The plateu can be inferred to be hydrogen-rich at the time of collapse by the presence of an extended plateauto deter phase in the SN light curve, and are classified as Type IIP. % \\citep{Li2011}. %The inferrence is made from the presence of H I lines in the spectrum, combined with an extended plateau phase, classifying these objects as Type IIP. %, and are believed to be collapsing stars with ZAMS progenitor masses in the $\\sim$8-30 \\msun~range that have retained most of their hydrogen envelopes throughout their evolution \\citep[e.g.][]{Heger2003}. %The plateau is produced by a recombination front that moves inward  in Lagrangian coordinates, but stays at roughly constant radius in the observer's frame \\citep{Grassberg1971}. After the plateau phase, the core of the SN becomes visible, glowing from radioactive input by \\iso{56}Co. As long as the ejecta remain opaque to the gamma-rays emitted in the decay, the light curve follows the exponential decay of \\iso{56}Co, with an $e$-folding time of 111.4 days. The spectrum evolves from being dominated by absorption lines and lines exhibiting P-Cygni profiles superimposed upon a blackbody-like continuum to strong emission lines and a weaker continuum. As the temperature falls, thermal emission eventually shifts into the infrared, but ultraviolet/optical features remain, caused by non-thermal ionizations and excitations.  %\\paragraph{METHODS TO DETERMINE THE PROGENITOR MASS}% --------------------- %\\textbf{Progenitor analysis}\\\\ From hydrodynamical modeling, the progenitors of Type IIP SNe are believed to be red supergiants (RSGs) \\citep{Chevalier1976, Falk1977}. %Until recently, however, there was no direct observational evidence for this. The rapidly growing image archives have improved the prospects of confirming this by identifying SN progenitors in imaging surveys. Such identifications have now been made for the Type IIP SNe SN 2003gd, SN 2004A, SN 2004et, SN 2005cs, SN 2008bk, SN 2008cn, and SN 2009md \\citep[][and references therein]{Smartt2009a, EliasRosa2009, Fraser2011}. For SN 2003gd, the progenitor identification was verified by follow-up observations showing that the candidate had disappeared \\citep{Maund2009}. In the cases where multiple-band detections were made, the progenitors have shown properties consistent with the RSG hypothesis, except for SN 2008cn, where the progenitor was yellow\\footnote{There is, however, the possibility that the detected source is a blend of two or more stars, and SN 2008cn differs from normal Type IIP SNe in its plateau length.} \\citep{EliasRosa2009}.   The estimated luminosity of the progenitor allows its ZAMS mass to be determined from stellar evolution models. The resulting values are consistently found to be below 20 \\msun, with a statistical analysis yielding a range of $8.5_{-1.5}^{+1}-16.5_{-1.5}^{+1.5}$ \\msun~for the progenitor population \\citep{Smartt2009b}, although dust effects may produce a somewhat higher upper boundary \\citep{Walmswell2012}. The fate of stars in the $20-25/30$ \\msun~range then remains unclear, as standard stellar evolution models predict also these to evolve to RSGs and explode as Type IIP SNe \\citep[e.g.,][]{HegerLanger2000, Meynet2003, Eldridge2008}. Models including rotation show that stars with initial mass greater than 20 M$_{\\odot}$ and time-averaged equatorial velocities of $\\gtrsim$200\\,km s$^{-1}$ could lose much of their hydrogen envelope and move bluewards in the HR diagram \\citep[e.g.][]{Meynet2003}. While this is immediately attractive as an explanation for the lack of luminous RSGs as Type IIP progenitors, one must account for the broad observed distribution of equatorial velocities \\citep[e.g.][]{Hunter2008} as well as the lack of detection of blue, luminous progenitors for Type IIn/Ib/Ic SNe \\citep{Smartt2009a}. %RSGs with $M_{\\rm bol}$ up to -9, corresponding to $M_{\\rm ZAMS}\\sim 30$ \\msun~stars, are indeed observed in the Milky Way and the Magellanic Clouds \\citep{Levesque2005, Levesque2006},  although one should be aware that estimating accurate bolometric magnitudes for such red stars is difficult, depending on model atmospheres for the bolometric corrections. Curiosuly however, all the RSGs in M31 correspond to $M_{\\rm ZAMS}\\lesssim 20$ \\msun.  %\\textbf{b) Photospheric phase modeling}\\\\ Another method to derive information about the exploded star is through radiation-hydrodynamical modeling of the bolometric light curve. \\citet{Litvinova1983, Lit1985} presented scaling relations from fits to a grid of Type IIP SN models allowing for the determination of the explosion energy, ejecta mass, and progenitor radius from the observed plateau length, $V$-band magnitude, and photospheric velocity. However, the neglect of \\iso{56}Ni in these models renders them viable only for SNe with very small \\iso{56}Ni masses. \\citet{Eastman1994} found that a (typical) \\iso{56}Ni mass of 0.06 \\msun~extends the plateau by 40\\%, and \\citet{Kasen2009} found a prolongation by 20-30\\% for \\iso{56}Ni-masses of $0.05-0.1$ \\msun. Since the ejecta mass has a strong scaling with the plateau length (to the power 2.9) in the \\citet{Lit1985} relations, a significant overestimate of the ejecta mass results. %Another issue with these scalings is that they are difficult to apply in practice, relying on $V$ and $u_{\\rm ph}$ at just one point on the lightcurve (the mid-plateau), and requireing an accurate determination of the photospheric velocity from the spectra (which needs to be taken at mid-plateau). %The rms for $M_{\\rm ej}$ in these fits is 32\\%. %\\citet{Hamuy2003} used these relat  More reliable results can be obtained by fitting hydrodynamical models including \\iso{56}Ni to the whole light curve evolution. %ions to determine the explosion parameters for 13 Type IIP SNe, obtaining ejecta masses in the range $17-56$ \\msun, and %\\citet{Nad2003} used these relations on a sample of 12 IIP SNe, finding $M_{\\rm ej}=10-32$ \\msun~with a median value of 15 \\msun, meaning that all objects would have $M_{\\rm ZAMS} \\gtrsim 12$ \\msun, and that $M_{\\rm ZAMS}^{\\rm median} \\gtrsim 17$ \\msun. Indeed t % obtaining ejecta masses $M_{\\rm ej}=10.3-30$ \\msun. %explosion energies $E=(0.5-3)\\e{51}~B$ (1 B = 1 Bethe =  $10^{51}~\\mbox{ergs}$), radii $R_0=200-600~R_\\odot$,~and ejecta masses $M_{\\rm ej}=10-30$ \\msun. %Radiation-hydrodynamical modeling of Such modeling has been undertaken for SN 1999em \\citep{Baklanov2005, Utrobin2007, Bersten2011}, SN 2003Z \\citep{UCP2007}, SN 2004et \\citep{Utrobin2009}, SN 2005cs \\citep{Utrobin2008}, and SN 2009kf \\citep{Utrobin2010}, obtaining %energies in the range $E=(0.25-2.3)\\e{51}~\\mbox{ergs}$, radii $R_0=230-1500~R_\\odot$, and ejecta masses $M_{\\rm ej}=14-28$ \\msun. %Modeling of individual SNe has also been formed, discussed in more detail in the discussion. As discussed by \\citet{Utrobin2009} and \\citet[][M10 hereafter]{Maguire2010}, these ejecta masses are generally too high to be consistent with the initial masses determined from direct observations of SN progenitors, as well as for what is expected from stellar evolution in general. Recently, \\citet{Inserra2011} and \\citet{Inserra2012}, using the radiation hydrodynamics code of \\citet{Pumo2011}, determined lower ejecta masses of $M_{\\rm ej}=5-7.5$ and $M_{\\rm ej}=8-12$ \\msun\\ for the Type IIP SNe 2007od and SN 2009bw.%, perhaps showing signs of alleviating this problem. %(a Salpeter IMF with a progenitor range of $8-30$ \\msun~would give some $M_{\\rm ej}\\sim$ 6 \\msun~and $M_{\\rm ZAMS}^{\\rm median}=12$ \\msun). Indeed, adopting the results from radiation-hydrodynamical modeling, the fate of stars in the $8-12$ $M_\\odot$ range remains unclear. Radiation-hydrodynamical modeling and direct observation of their progenitors do, however, agree on the progenitors of IIP SNe being RSGs.  A third method for diagnosing the progenitor mass is through late-time spectral modeling. In the nebular phase ($t \\gtrsim 150$ days), the inner ejecta become visible, and the various nuclear burning zones can be analysed. Stellar evolution models predict the metal core mass to strongly increase with progenitor ZAMS mass \\citep[e.g.][]{WW95}, making it possible to distinguish between different progenitors by determining the nucleosynthesis yields. Emission lines of carbon, oxygen, neon, sodium, magnesium, silicon, and sulphur are the main signatures that can potentially constrain the progenitor mass.%amount of metals produced.  %As pointed out by \\citet{Dessart2010}, the Type IIP class guarantees that no core material has been removed by mass loss, which is not the case for Ic's, where the progenitor's ZAMS mass and its final core mass cannot be linked.  Despite being the most common core-collapse SN class, there are so far only a handful of Type IIP SNe that have extensive temporal and spectral coverage in the nebular phase, and not much information has as yet been derived from these observations. The best existing datasets are for SN 1990E \\citep{Schmidt1993, Gomez1993, Benetti1994}, SN 1999em \\citep{Leonard2002,Elmhamdi2003}, SN 1999gi \\citep{Leonard2002gi}, SN 2002hh \\citep{Pozzo2006}, SN 2004dj \\citep{Chugai2005}, SN 2004et \\citep[][M10]{Sahu2006}, SN 2005cs \\citep[e.g.][]{Pastorello2009}, 2007it \\citep{Andrews2011}, and 2007od \\citep{Inserra2011}. Of these, only SN 1999em was well monitored also in the near-infrared (NIR). % (Hamuy2001, Elmhamdhi2003b,Krisciunas2009). 2004dj : 7 optiska spectra up till 350 dagar (Chugai2005 %Since the mass of the helium core, and the various zones therein, depends sensitively on the progenitor mass (e.g. WoosleyWeaver1995), this has potential for being a sensitive technique. the  Apart from the sparsity of data, spectral analysis needs complex tools to model the transformation of energy from radioactive decay to emergent UVOIR radiation. Temperature and NLTE solutions for a large number of zones and elements, including non-thermal heating/ionization/excitation rates, and multi-line radiative transfer are the fundamental ingredients that must be calculated, together making up millions of constraints whose self-consistent solution must be iteratively sought. In addition, molecules, dust, and time-dependent effects may cause further complications.  So far, SN 1987A is the only Type II SN for which detailed spectral modeling has been undertaken.  Despite being a Type IIpec rather than a Type IIP, it was probably similar to a Type IIP in the nebular phase, since the nucleosynthesis is largely unaffected by the late-time evolution of the envelope. \\citet{Xu1991} and \\citet{Kozma1992} obtained the solutions for the ionization, excitation, and heating produced by the gamma-rays and positrons, which is the first step in the nebular phase modeling. The evolution of individual lines were analyzed in a series of papers by \\citet{Xu1992} (hydrogen), \\citet{Li1992} (oxygen), \\citet{Li1993iron} (iron, cobalt, nickel), \\citet{Li1993} (calcium), and \\citet{Li1995} (helium). The thermal evolution of the envelope was modelled in detail by \\citet{deKool1998}, and \\citet[][in the following KF98a,b]{Kozma1998I, Kozma1998II} computed the spectra from detailed explosion models to study the evolution of temperature, ionization, and line fluxes in the $200-800$ day range. \\citet{Kjaer2010} and \\citet[][J11 hereafter]{Jerkstrand2011} analyzed the spectrum in the \\iso{44}Ti-powered phase ($t \\gtrsim 5$ years), including the effects of multi-line radiative transfer.%They estimate a total H-zone mass of 7.7 \\msun, of which 2.2 \\msun~resides the core ($V<2500~$\\kms). The oxygen mass of $\\sim$ 1.5 \\msun~in the explosion models is found to reproduce the [O I]\\wll6300, 6364 lines well. They also make the important point that many lines originate in other zones than where most of the element mass resides. Many of the Fe II lines are found to have strong contributions by primordial iron in the hydrogen zone at all times.  Modeling of other objects includes the work by \\citet{Dessart2011}, who compared the emergent spectra of Type IIP explosion models to SN 1999em, assuming complete thermalization of the gamma-rays, but performing a detailed radiative transfer calculation. %They show that higher mass progenitors produce stronger and broader [O I]\\wll6300, 6364 lines. They also find that the 8498 and 8542 \\AA~Ca II photons scatter in the 8662 line, resulting in a very broad 8662 \\AA~component. The [O I]\\wll 6300, 6364 doublet has a significant contamination by iron lines at late times. %We comment on their modeling results for Type IIP SNe compared to ours in the discussion.  In this paper we undertake a detailed analysis of  one of the brightest and best-observed Type IIP SNe to date, SN 2004et. This SN has been subject to progenitor analysis, hydrodynamical modeling, and some qualitative spectral analysis. \\citet{Li2005} identified a progenitor candidate in ground-based pre-explosion images. However, \\citet {Crockett2011} showed that this candidate was still visible in 2007 (3 years after explosion), and that the source in the original images was a composite of two or three sources. % in the ground-based imaging. %Furthermore, the progenitor candidate was resolved into three different components of which one, called object Centre, is coincident with the SN. The analysis results in two scenarios differing in how much flux the SN is assumed to contribute to object Centre. In the first scenario, object Centre is assumed to consist solely of the SN, and in the second scenario, it is assumed to consist partly of SN light and partly of another source. In both cases They identified an excess flux in the pre-explosion images contributed by the true SN progenitor, and for two different assumptions about the late-time SN flux determined $M_{\\rm ZAMS}=8$ and 10 \\msun, with an error of $+5$/$-1$ \\msun. With new post-explosion imaging and updated methods to determine bolometric corrections and synthetic colours from model spectra, Fraser et al. 2012 (in prep.) have revised this to $M_{\\rm ZAMS}=11_{-1}^{+2}$ \\msun. The best estimates for luminosity, temperature, and radius are log $\\left(L/L_\\odot\\right) = 4.8$, $T= 3600$ K, and $R_0 =  650 ~R_\\odot$.  Through radiation-hydrodynamical modeling, \\cite{Utrobin2009} derived a large ejecta mass of $M_{\\rm ej}=24.5\\pm 1~M_\\odot$, the highest value for any Type IIP modeled in detail so far. They also determined a progenitor radius $R_0=1500\\pm 140~R_\\odot$, explosion energy $E=2.3\\pm 0.3~\\mbox{B}$ (1 B = $10^{51}$ ergs) and \\iso{56}Ni mass $M(^{56}\\mbox{Ni})=0.068\\pm 0.009 M_\\odot$.  Spectroscopic analysis is so far limited to some qualitative results from the nebular spectrum by \\citet[][S06 hereafter]{Sahu2006} and M10. S06 found the [\\ion{O}{i}] \\wll6300, 6364 luminosity to be of comparable to the one in SN 1987A, and from the similar \\iso{56}Ni mass they conclude that SN 2004et should have a similar oxygen mass. From this follows a similar progenitor mass, which was $16-22$ \\msun~ for SN 1987A \\citep{Arnett1989}. %WH07 : 18-20. %and a comparison of the O I\\wll6300,6364 lines to 1987A suggested just a somewhat smaller mass than 87A had ($\\sim$19\\msun). M10 found the [\\ion{O}{i}] \\wll6300, 6364 flux normalized to the \\iso{56}Ni mass to be 15\\% lower than in SN 1987A at 285 days, and assuming all other things equal the oxygen mass would then be smaller by the same factor. With the strong dependency of oxygen production on progenitor mass \\citep[e.g.][]{WW95}, SN 2004et would then be expected to have a progenitor mass only slightly lower than SN 1987A.  Taken together with the progenitor analysis and light curve modeling, there are therefore conflicting results for how massive the progenitor of SN 2004et was. The aim in this paper is to shed more light on this issue by modeling the photometric and spectroscopic nebular phase data. To this end, we use the spectral synthesis code described in J11 (see also \\citet{Maurer2011} for testing of the code), with some modifications described in Appendix \\ref{sec:updates}. We compare the model spectra produced by using three different explosion models as input  (from $M_{\\rm ZAMS}=$ 12, 15 and 19 \\msun\\ progenitors), and investigate which produces best agreement with observations. This paper is complemented by a more general discussion of the nebular phase spectra of Type IIP SNe in \\citet[][in the following M12]{Maguire2012}. ",
        "conclusions": "\\begin{itemize} \\item We have shown that nebular-phase spectral modeling can be used to constrain the progenitor masses of Type IIP supernovae. Lines showing significant dependency on the progenitor mass are [\\ion{O}{i}] \\wll6300, 6364,  \\ion{Na}{i} \\wll5890, 5896, \\ion{Mg}{i}] \\wl4571, \\ion{Mg}{i} \\wl1.503 $\\mu$m, \\ion{C}{i} \\wl1.176 $\\mu$m, \\ion{C}{i} \\wl1.454 $\\mu$m, \\ion{Si}{i} \\wl1.20 $\\mu$m, [\\ion{Si}{i}] \\wll1.607, 1.645 $\\mu$m, and [\\ion{Ne}{ii}] \\wl12.81 $\\mu$m. \\item The nebular-phase optical and near-infrared emission lines of SN 2004et are well reproduced by model spectra of a $M_{\\rm ZAMS}=15$ \\msun~(non-rotating) progenitor star, with an oxygen mass of 0.8 \\msun. Spectra from 12 and 19 \\msun~progenitors under- and overproduce most of the nebular phase lines, respectively, with particular discrepancy for the high-mass model. This result is in fair agreement with the analysis of pre-explosion imaging ($M_{\\rm ZAMS}= 11_{-1}^{+2} M_\\odot$, Fraser et al. 2012, in prep.), but in disagreement with the progenitor mass inferred from hydrodynamical modeling ($M_{\\rm ZAMS}> $25 \\msun, \\citet{Utrobin2009}). \\item In the $300-600$ day range, we find that the [\\ion{O}{i}] 6300, 6364 doublet emits 1, 2, and $4-6$\\% of the \\iso{56}Co decay energy for 12, 15, and 19 \\msun\\ progenitor stars. This result may be used for a coarse diagnosis of other Type IIP SNe with similar \\iso{56}Ni mass (0.06 \\msun). \\item By modeling the mid-infrared iron-group lines observed with Spitzer, we determine a filling factor $f=0.15$ for the Fe/Co/Ni clumps in the ejecta, corresponding to a density $\\rho(t) = 7\\e{-14}\\times (t/100\\ \\mbox{d})^{-3}$ g cm$^{-3}$. This is similar to the filling factor derived for SN 1987A by \\citet{Kozma1998II}, whereas \\citet{Li1993iron} likely overestimate the filling factor due to neglect of emission by primordial iron in the hydrogen zone. \\item By our calculation of the atomic emission lines and continua, we can confirm the contribution by a silicate dust feature to the mid-infrared spectrum, first identified by \\citet{Kotak2009}, as well as the presence of SiO and CO fundamental band emission in the spectrum. The CO luminosity matches the total heating of the O/C zone in a $12-15$ \\msun~progenitor.  \\end{itemize}"
    },
    "1208/1208.0002_arXiv.txt": {
        "abstract": "We introduce the Making Galaxies in a Cosmological Context (MaGICC) program of smoothed particle hydrodynamics (SPH) simulations.  We describe a parameter study of galaxy formation simulations of an L$^\\star$ galaxy that uses early stellar feedback combined with supernova feedback to match the stellar mass--halo mass relationship.  While supernova feedback alone can reduce star formation enough to match the stellar mass--halo mass relationship, the galaxy forms too many stars before $z=2$ to match the evolution seen using abundance matching.  Our early stellar feedback is purely thermal and thus operates like a UV ionization source as well as providing some additional pressure from the radiation of massive, young stars.  The early feedback heats gas to $>10^6$ K before cooling to $10^4$ K.  The pressure from this hot gas creates a more extended disk and prevents more star formation prior to $z=1$ than supernovae feedback alone.   The resulting disk galaxy has a flat rotation curve, an exponential surface brightness profile, and matches a wide range of disk scaling relationships.  The disk forms from the inside-out with an increasing exponential scale length as the galaxy evolves.  Overall, early stellar feedback helps to simulate galaxies that match observational results at low and high redshifts. ",
        "introduction": "As matter collapses in the early Universe, gas initially heats as kinetic energy is converted into thermal \\citep{Rees1977, Birnboim2003}.  In massive halos, the gas reaches temperatures where its cooling time becomes longer than the Hubble time \\citep{Rees1977}, which results in a theoretical maximum galaxy mass.  However, in the centers of galaxies, enough gas accumulates to increase the density to the point that the gas efficiently radiates and cools.  Such cooling is unstable since the radiation leaves through the optically thin surrounding hot gas.  This process leads to overcooling since as the hot gas cools, it stops providing pressure support for the surrounding gas, more dense gas is able to accumulate and cool even more efficiently \\citep{White1978}.  This problem is known as the ``overcooling catastrophe'' \\citep[e.g.][]{Balogh2001}.  The ``overcooling catastrophe'' has long plagued simulations of disk galaxy formation.  Fully cosmological numerical galaxy simulations consistently contain a massive central concentration of stars \\citep{Navarro1991, Governato2004, Stinson2010, Scannapieco2012}.  This is evident in the central peak in rotation curves in simulated galaxies as well as in their surface brightness profiles.  Further evidence for simulations forming too many stars comes from recent studies matching observations of galaxy stellar masses with their dark matter halo masses.  One method is the abundance matching technique.  To lay the groundwork for this technique, \\citet{Conroy2006} rank ordered halos by total halo mass from collisionless simulations and then did a rank ordering of galaxies based on their stellar mass from the Sloan Digital Sky Survey.  The halos and galaxies were then divided in several mass bins.  \\citet{Conroy2006} found that the correlation functions between halos in certain mass bins were well matched by the correlation functions of galaxies in corresponding stellar mass bins.  Several groups used this abundance matching technique to compare total halo masses with galaxy stellar masses \\citep{Conroy2009,Moster2010,Guo2010,Behroozi2010}.  Direct comparisons of stellar mass and total halo mass have also been done based on satellite dynamics \\citep{Klypin2009, More2009, More2011} and weak lensing \\citep{Mandelbaum2009, Schulz2010}.  For a comparison of all these techniques, please see Figure 11 of \\citet{Behroozi2010}.  Each method shows a good level of correspondence.  \\citet{Guo2010} and \\citet{Sawala2011} showed that nearly all the simulations of galaxy formation have formed many more stars than abundance matching predicts.  Even without overcooling, \\citet{vandenBosch2002} showed that the amount of low angular momentum material in collapsed collisionless halos exceeds the amount of low angular momentum material observed in disk galaxies.  Consequently, this low angular momentum material needs to be removed from the center of the system.  Stellar feedback is the favored way of reducing  star formation and launching outflows \\citep{Scannapieco2008, Schaye2010}.  \\citet{Dekel1986} showed that supernova feedback can eject gas from galaxies with virial velocities up to 100 km s$^{-1}$.  Semi-analytic models have found that significant amounts of stellar feedback are required to match the low mass end of the luminosity function \\citep{Somerville1999, Benson2003, Bower2006, Bower2008, Bower2012}. \\citet{Dutton2009} showed that stellar feedback can remove low angular momentum material.  In hydrodynamical simulations, two methods are commonly used to model stellar feedback.  One is kinetic feedback that adds velocity kicks to gas particles to remove them from the inner regions of galaxy disks \\citep{SH03, Oppenheimer2006, DallaVecchia2008}.  The other is thermal feedback in which stars simply heat gas particles and allow the adiabatic work of the particles to push other gas out of the way \\citep{Gerritsen1997, Thacker2000, Kawata2003, Stinson2006}.  Since stars form in dense regions, the cooling times of the surrounding gas are short, and without help, the gas will quickly radiate away all the supernova energy \\citep{Katz1992}.  In real galaxies, the amount of gas necessary to exert a dynamical influence on the ISM small.  In simulations, such small amounts of gas are difficult to model, so a common technique has been to turn off cooling for a limited amount of time \\citep{Gerritsen1997, Thacker2000, Brook2004, Stinson2006}.  It is not clear that kinetic and thermal feedback have significantly different effects than one another.  \\citet{Durier2012} showed that kinetic feedback has the same effect as thermal feedback when the hydrodynamics is left turned on.  However, in many implementations of kinetic feedback, the hydrodynamic processes are disabled for a set period of time to maintain numerical convergence \\citep{SH03, Oppenheimer2006}.  It is also possible that thermal feedback does not need to rely on disabling cooling.  \\citet{DallaVecchia2012} recently showed that thermal feedback can have a significant dynamical effect if the total energy of all the supernovae explosions is injected into one particle at one time.  This large, one time energy deposition raises the particle's temperature high enough that its cooling time lengthens enough for the feedback to have a dynamical effect.  For the first time in simulations, \\citet{Governato2010} showed that thermal stellar feedback can remove low angular momentum material from dwarf galaxies, creating galaxies with slowly rising rotation curves.  \\citet{Brook2011} showed that the ejected low angular momentum material becomes part of a galactic fountain and can be reaccreted onto the disk with higher angular momentum than it left. \\citet{Sawala2011} showed that even using a high density threshold for star formation that made the stellar feedback more efficient, these simulations produced more stars than is predicted by the stellar mass--halo mass relationship.  Three other recent simulations have shown success at forming realistic $L^\\star$ disk galaxies.  \\citet{Guedes2011} used high resolution and a high amount of supernova feedback to produce a realistic disk galaxy.  \\citet{Agertz2011} suggested that using a low star forming efficiency is helpful for creating extended disks.  They also use models that increase the amount of supernova energy deposited above the canonical $10^{51}$ erg per supernova.  These simulations do the best job flattening out the central region of their rotation curve.  \\citet{McCarthy2012} also form galaxies with flat rotation curves in a large, low resolution cosmological volume.  They employ a kinetic feedback scheme that uses less than the canonical $10^{51}$ erg supernova energy, though see \\citet{Kay2002} for a demonstration of how kinetic feedback provides a larger effect than thermal feedback.  The galaxies each contain a few thousand particles and employ a gravitational softening of $\\sim1$ kpc, which makes an examination of the detailed structure of the galaxies challenging.  These simulations show that it is possible that more feedback than the canonical $10^{51}$ erg is required.  Instead of simply increasing the amount of energy released by supernovae feedback, we note that \\citet{Murray2010} showed that there can be significant feedback effects from stars before they explode as supernovae.  \\citet{Hopkins2011} incorporated a kinetic radiation pressure feedback into simulations of isolated disk galaxies and found that the feedback could strongly regulate star formation.  While we do not have the resolution in cosmological simulations to implement a similar kinetic scheme, we implement a scheme based on thermal pressure to provide feedback during the time between when stars are formed and the first SN star exploding.  Using this feedback prescription, \\citet{Maccio2012} showed that the feedback removes low angular momentum dark matter in galaxies up to nearly L$^\\star$, producing cored dark matter density profiles. \\citet{Stinson2012} showed that the metal rich outflows created by this feedback match observations of OVI in the circum-galactic medium of star forming galaxies.  \\citeauthor{Brook2012a} (2012b) also showed that a sample of lower mass galaxies form disks that follow a wide range of disk scaling relationships.  Here, we present a detailed study of how varying the key free parameters in our simulations affect the morphology and evolution of galaxies.  In \\S \\ref{sec:sims}, we describe the galaxy we model and the physics used in the simulation.  In \\S \\ref{sec:results}, we show how varying parameters affects the mass of stars formed, the morphology of the galaxy, and how the galaxy evolved. ",
        "conclusions": "We present a parameter study that varies the strength of stellar feedback in simulations of galaxy formation.  As a constraint, we use the stellar mass--halo mass relationship.  To reduce star formation enough to match the relationship, early stellar feedback from massive stars was required before they explode as supernova.  Our fiducial model best fits the stellar mass--halo mass relationship, the evolution of that relationship, has a flat rotation curve and an exponential surface brightness profile with a modest bulge in the center.  In contrast to many previous cosmological galaxy formation simulations, most of the star formation occurs after $z=2$.  There are two peaks of star formation that correspond to the two significant mergers to the galaxy.  Both of these mergers are gas rich.  We find that the mass of stars formed at $z=0$ is very sensitive to the amount and timing of stellar feedback employed.  Models that form too many stars follow a common pattern where they turn half of their baryons into stars, the other half is hot gas.  These galaxies also form many of their stars at higher redshift than what \\citet{Moster2012} find when they compare high redshift luminosity functions with N-body simulations.  Galaxies that form too many stars have bright central concentrations of stars.  This is reflected in galaxy rotation curves as a high central peak.  Galaxies that form the right amount of stars have exponential surface brightness profiles and slowly rising rotation curves.  A simulation which does not use early stellar feedback, but increases the supernova energy to $1.2\\times10^{51}$ erg ends up with a similar stellar mass to the fiducial simulation, but does not compare as well with other observed galaxy properties.  The difference is that our implementation of early stellar feedback keeps the gas above $10^4$ K and at densities below 10 cm$^{-3}$, an effect similar to that achieved by UV ionization.  The low density, warm gas keeps the disk extended prior to $z=1$ and thus keeps gas out of the central 2 kpc of the galaxy.  Since less gas makes it to the central region, fewer stars form before $z=1$ and, therefore, this model does not result in the massive central concentration that forms without early stellar feedback.  The early stellar feedback produces galaxies that correspond to observations in a number of ways.  Using early stellar feedback in dwarf galaxies spanning a wide mass range, \\citeauthor{Brook2012a} (2012b) showed that the simulated galaxies match the disk scale length, gas fractions, and luminosities of observed galaxies.  \\citeauthor{Brook2012} (2012a) has shown that the reason for such agreement is that outflows redistribute low angular momentum gas from the centers of galaxies to their outskirts.  \\citet{Stinson2012} showed that the outflows that redistribute the gas also populate the circum-galactic medium with an amount of oxygen that closely matches observations of OVI in the CGM.  \\citet{Maccio2012} showed that these outflows can also change the inner regions of dark matter density profiles from steeply rising cusps into flat cores.  In Kannan et al (\\emph{in prep}), we will show how the stellar feedback described here performs in a larger sample of lower resolution galaxies in a cosmological volume.  Early stellar feedback helps keeps gas out of the center of galaxies, which leads to forming disk galaxies like those that we observe."
    },
    "1208/1208.0833_arXiv.txt": {
        "abstract": "{In a recent paper arXiv:1107.5048, we discussed the correlation between the elastic neutralino-nucleon scattering cross section, constrained by dark matter direct detection experiments, and fine-tuning at tree-level in the electroweak symmetry breaking sector of the Minimal Supersymmetric Standard Model (MSSM). Here, we show that the correlation persists in the Next-to-Minimal Supersymmetric Standard Model (NMSSM), and its variant, $\\lambda$-SUSY. Both models are strongly motivated by the recent discovery of a 125 GeV Higgs-like particle. We also discuss the implications of the recently published bound on the direct detection cross section from 225 live days of XENON100 experiment. In both the MSSM and the NMSSM, most of the parameter space with fine-tuning less than 10\\% is inconsistent with the XENON100 bound. In $\\lambda$-SUSY, on the other hand, large regions of completely natural electroweak symmetry breaking are still allowed, primarily due to a parametric suppression of fine-tuning with large $\\lambda$. The upcoming XENON1T experiment will be able to probe most of the parameter space with less than 1\\% fine-tuning in all three models. } ",
        "introduction": "Several experiments around the world are currently attempting to observe dark matter via ``direct detection\", {\\it i.e.} measuring recoils of detector nuclei following their collisions with ambient dark matter particles. While no convincing observation has been reported so far, the sensitivity of the experiments is rapidly increasing. Currently, the best upper bounds on the cross section of elastic, spin-independent dark matter-nucleon scattering in the 10 GeV-TeV mass range come from the XENON100 experiment~\\cite{xenon100paper,xenon2012}, and are of the order $10^{-44}-10^{-45}$ cm$^2$. This is the range where the predictions of many attractive theoretical models of dark matter lie~\\cite{DMreviews}. The most studied of these is R-parity conserving supersymmetry, where the lightest neutralino $\\chi^0_1$ generically has the right properties to explain dark matter. It is therefore important to understand the implications of the direct detection bounds for supersymmetric dark matter.  In a recent paper~\\cite{ftmssm}, we pointed out a strong correlation between direct detection cross sections and naturalness of electroweak symmetry breaking (EWSB) in the Minimal Supersymmetric Standard Model (MSSM)\\footnote{Related discussions have also appeared in Refs.~\\cite{others,zeptobarn}.}: MSSM parameter points with lower direct detection cross section have more finely-tuned EWSB. This conclusion seems very general: It does not depend on the details of the SUSY-breaking mechanism and high-scale physics, and for most of the parameter space (except pure-Higgsino dark matter) it does not require imposing the thermal relic density constraint, so that it applies even in models with non-standard early cosmological history. The only assumption is the absence of accidental cancellations among physically distinct contributions to the dark matter-nucleon scattering amplitudes ({\\it e.g.} Higgs- and squark-exchange diagrams).  The main goal of the present paper is to extend this study to the Next-to-Minimal Supersymmetric Standard Model (NMSSM). (For reviews, see, for example, Refs.~\\cite{mainref,NMSSMreview}.). Our main motivation to study the NMSSM is the recent discovery of a new particle, with a mass of about 125 GeV and properties consistent with the Standard Model (SM) Higgs, at the LHC~\\cite{LHC_Higgs}. In the MSSM, a Higgs particle of this mass can only be accommodated at a price of severe fine-tuning in the EWSB~\\cite{HPR}. Fine-tuning is significantly alleviated if an extra singlet (with respect to SM gauge groups) superfield is coupled to the Higgs sector, as in the NMSSM~\\cite{HPR}. In this respect, a particularly promising variation of the NMSSM is the ``$\\lambda$-SUSY\"~\\cite{lambdasusy},  characterized by a large superpotential coupling $\\lambda$ between the singlet and doublet Higgs fields. The addition of a singlet has a non-trivial effect on dark matter phenomenology due to the possible admixture of the singlino in $\\chi^0_1$, as well as the additional SM-singlet Higgs state. Does the correlation between direct-detection rates and EWSB fine-tuning persist in the NMSSM and $\\lambda$-SUSY? It should be noted that in \\cite{ftmssm} and in this paper, we are only interested in the tree-level naturalness of the $Z$ mass. Of course, additional fine-tuning may be induced by the radiative corrections associated with heavy stops and gluinos; in this sense, the fine-tuning measure we use should be interpreted as the lower bound on the total fine-tuning. We are not aware of any correlation between the fine-tuning due to radiative corrections and dark matter direct detection cross sections.  This paper is organized as follows. We begin by reviewing and updating the MSSM results of Ref.~\\cite{ftmssm} in Section~\\ref{sec:review}, including the latest results from XENON100~\\cite{xenon2012}. We then briefly review the structure of the NMSSM and $\\lambda$-SUSY in Section~\\ref{sec:NMSSM}. Section~\\ref{sec:FT} provides definitions of fine-tuning, and discusses a suppression of fine-tuning at large values of $\\lambda$, which will be important for understanding our results. The analysis procedure is outlined in Section~\\ref{sec:scan}, and the results are presented and discussed in Section~\\ref{sec:results}. Finally, we conclude in Section~\\ref{sec:conc} with a brief summary of the main results and directions for future work. ",
        "conclusions": "\\label{sec:conc}  In this paper, we continued our study, initiated in Ref.~\\cite{ftmssm}, of the implications of dark matter direct detection experiments on supersymmetric models. In Ref.~\\cite{ftmssm} we found that the MSSM, for ``generic\" parameters, predicts a spin-independent elastic neutralino-nucleon scattering cross section of the order of $10^{-45}-10^{-44}$ cm$^2$ (depending somewhat on the poorly known strange quark form factor) or higher. Cross sections in this range are currently being probed by XENON100, which already places meaningful constraints. Suppressing the cross section below the ``generic\" level in the MSSM requires (barring accidental cancellations) that the LSP be either a pure gaugino or a pure Higgsino. In the first case, lowering the direct detection cross section requires raising $\\mu$, and therefore introducing fine-tuning in the EWSB. In the second case, requiring that the Higgsino be a thermal relic implies EWSB fine-tuning of about 1/500, independent of direct detection bounds.  The recent LHC discovery of a new particle with 125 GeV mass and properties consistent with the SM Higgs puts significant pressure on the MSSM, since fine-tuning of order $0.1$\\% is required to accommodate it in this model. This motivates considering supersymmetric models with non-minimal Higgs sectors, where new contributions to the tree-level Higgs mass can easily arise. The simplest example is the NMSSM, where a single gauge-singlet superfield is added, and a 125 GeV Higgs can be incorporated with far less fine-tuning. The reduction of tuning is especially striking in the version of the NMSSM with strong doublet-singlet Higgs coupling, the so-called $\\lambda$-SUSY. In this paper, we extended the analysis of Ref.~\\cite{ftmssm} to the NMSSM and $\\lambda$-SUSY. We found that the qualitative correlations between the dark matter direct detection cross section, the LSP composition, and the EWSB fine-tuning found in the MSSM essentially persist in these non-minimal models as well. Numerically, the minimal cross section allowed for the same level of EWSB fine-tuning is somewhat decreased in the NMSSM compared to the MSSM, and is further decreased, rather significantly, in $\\lambda$-SUSY. We discussed the physical origin of these effects. We found that the current XENON100 cross section bounds are in mild tension with the MSSM and the NMSSM, excluding most points with fine-tuning of 1/10 or better, while large parts of completely natural parameter space are still allowed in $\\lambda$-SUSY.  As in Ref.~\\cite{ftmssm}, we took a deliberately broad approach to the supersymmetric model parameter space. We do not assume any particular model of SUSY breaking; instead, we scan over unconstrained weak-scale parameters of each model. We do not impose the relic density constraints on the LSP.\\footnote{Relic density constraints in the NMSSM are well known; see, for example, the review article~\\cite{mainref}. Relic density constraints in $\\lambda$-SUSY were considered in Ref.~\\cite{Cao:2008un}.}  We also do not fully utilize the recent data concerning the 125 GeV Higgs candidate: We only require broad consistency of the Higgs spectrum with the data, demanding that a tree-level mass of at least one mostly-doublet CP-even Higgs be between 100 and 150 GeV. Imposing any combination of additional constraints would select a subspace of the broad parameter space studied here. Stronger conclusions can be obtained with such added constraints, but they would be less generally applicable. It would be interesting to perform such studies in the future.  Looking ahead, we can anticipate further dramatic improvement in the experimental sensitivity of direct detection dark matter searches within a few years. Our results indicate that, if dark matter is supersymmetric, the searches will very likely be successful: Only supersymmetric models with sub-percent levels of EWSB fine-tuning, or accidental cancellations, will escape detection by experiments with sensitivity levels expected of, for example, XENON1T. In this paper, we showed that these statements apply not only in the MSSM but also in well-motivated non-minimal supersymmetric models. This underscores the importance of the continuing direct dark matter searches for fundamental physics.  \\vskip0.8cm \\noindent{\\large \\bf Acknowledgments} \\vskip0.3cm  This research is supported by the U.S. National Science Foundation through grant PHY-0757868 and CAREER grant PHY-0844667. MP would like to acknowledge the hospitality of the Aspen Center for Physics, supported by the NSF Grant \\#1066293. We would like to thank Roberto Franceschini, Aaron Pierce and David Sanford for useful discussions.  \\begin{appendix}"
    },
    "1208/1208.2231_arXiv.txt": {
        "abstract": "{Test particles infalling from infinity onto a compact spherical star with a mildly super-Eddington luminosity at its surface are typically trapped on the ``Eddington Capture Sphere\" and do not reach the surface of the star. The presence of a sphere on which radiation pressure balances gravity for static particles was first discovered some twenty five years ago. Subsequently, it was shown to be a capture sphere for particles in radial motion, and more recently also for particles in non-radial motion, in which the Poynting-Robertson radiation drag efficiently removes the orbital angular momentum of the particles, reducing it to zero. Here we develop this idea further, showing that ``levitation\" on the Eddington sphere (above the stellar surface) is a state of stable equilibrium, and discuss its implications for Hoyle-Lyttleton accretion onto a luminous star. When the Eddington sphere is present, the cross-section of a compact star for actual accretion is typically less than the geometrical cross-section $\\pi R^2$, direct infall onto the stellar surface only being possible for relativistic particles, with the required minimum particle velocity at infinity typically $\\sim1/2$ the speed of light.  We further show that particles on typical trajectories in the vicinity of the stellar surface will also be trapped on the Eddington Capture Sphere.} ",
        "introduction": "In Newtonian theory both gravity and radiative flux {\\it in vacuo} diminish like $1/r^2$ with the distance from the center of the star. The luminosity in units of the Eddington luminosity, $L(r)/L_{\\rm Edd}$, at {\\it any} distance from the star is equal to its value at the surface, $L(R)/L_{\\rm Edd}$, and is therefore {\\it  everywhere} either sub-Eddington, Eddington, or super-Eddington. However, in Einstein's general relativity the radiative force diminishes more strongly with distance than the gravitational force, the redshifted luminosity $L(r)$ decreasing as $1/(1-2R_G/r)$, and a static balance of forces is achieved at that radius $r$ at which $L(r)=L_{\\rm Edd}(1-~2R_G/r)^{-1/2}$. For this reason, radiation may be super-Eddington close to the star, but sub-Eddington further away, reaching the Eddington value at the ``Eddington Capture Sphere'', whose radius is (\\cite{Phinney}) \\begin{equation} \\label{Eddington-radius} R_{\\rm Edd} =  \\frac{2 R_G}{1 - \\left(1 - \\dfrac{2 R_G}{R} \\right)^2 \\left(\\dfrac{L(R)}{L_{\\rm Edd}} \\right)^2}\\ . \\end{equation} At this radius the gravitational and radiative forces balance (for a~static, optically thin hydrogen shell). The formula is in Schwarzschild coordinates, with the gravitational radius $R_G$ defined in Eq.~(\\ref{dimensionless-scaling}). \\cite{Abramowicz} (hereafter AEL) have rigorously shown by analytic calculations that in the combined gravitational and isotropic radiation fields of a~spherical, non-rotating, compact star,  radially moving test particles (with proton mass, and Thomson cross-section for photon momentum absorption) are captured on this sphere and ``levitate,'' i.e., remain at rest.  \\cite{Oh2011} have shown that when the  luminosity at the surface of the star is  mildly super-Eddington, particles with non-zero orbital angular momentum can  also be captured at $R_{\\rm Edd}$ (see also \\cite{Bini}). This is because the  Poynting-Robertson radiation drag acts as an effective torque,  and for a wide class of trajectories reduces the orbital  angular momentum to zero.  In this paper we further discuss the phenomenon of the Eddington Capture Sphere by considering two particular cases, one corresponding to initially quasi-circular motion close to the star, and the second being rather similar to the classic \\cite{Hoyle} accretion model. We also consider the question of stability and speculate about possible astrophysical manifestations of the Eddington Capture Sphere.  All numerical results presented here were obtained with the Dormand-Prince method, which is a fourth-order accuracy, adaptive step-size Runge-Kutta type integration method. ",
        "conclusions": "We have shown that a luminous star in Schwarzschild metric can capture test particles from a wide class of orbits onto a spherical surface of radius larger than the star itself.  The particles come to rest on this sphere, because radiation drag eventually removes all of their angular momentum, as shown by Oh, Kim, \\& Lee (2011). The particles remain suspended above the surface of the star, because every point on the capture sphere corresponds to a position of equilibrium, the equilibrium being stable in the radial direction (Abramowicz, Ellis, Lanza 1990) and neutral in directions tangent to the sphere. The radius of this ``Eddington Capture Sphere\" depends on the ratio of the stellar luminosity to the Eddington luminosity (Eq. [1]).  Thus, accreting compact stars, as well as stars which eject matter, may be surrounded by a shell of matter at rest, as long as the luminosity of the star is close to the Eddington luminosity (super-Eddington on the stellar surface).  The trajectories of particles falling from infinity in the gravitational field of a mildly super-Eddington compact star were investigated in Section 6. We showed that radiation drag strongly reduces both the azimuthal and the radial component of test particle velocities. This leads to a moderate enhancement of the capture cross-section by the Eddington sphere, relative to the cross-section for accretion by a non-radiating star, and a drastic suppression of the cross-section for actual accretion onto the stellar surface. Radiation drag also leads to a dramatic decrease in the kinetic energy deposited at the surface.  It is interesting to speculate that the presence of the Eddington Capture Sphere, an effect of Einstein's general relativity, may play a role in the formation and ejection of shells of matter by massive stars in the final stages of their evolution, such as Wolf-Rayet stars, or the luminous blue variables (LBV stars).  In principle, the capture sphere can exist at any radius outside the star, if the stellar luminosity has the appropriate value. However, the larger the radius of the sphere (in units of the Schwarzschild radius), the more finely tuned the luminosity has to be---for a $10R_G$ (5~Schwarzschild radii) capture sphere the luminosity at infinity has to be $0.95 L_{\\rm Edd}$, while for a 500 000 $R_G$ sphere the luminosity has to be equal to the Eddington value to an accuracy better than one part in a million.  While our discussion was couched in terms of the Thomson scattering cross-section, similar results will hold if other scattering or absorption processes contribute to momentum transfer from the radiation field to the particles. The capture sphere is located at that radius at which radiation pressure on matter at rest balances the gravitational pull of the star. Outside the sphere, the radiation pressure is too weak to balance gravity, inside the sphere it overcomes the gravitational pull.  It is expected that a more likely application of the Eddington capture sphere will be found in the interpretation of the behaviour of accreting neutron stars in low mass X-ray binaries. In particular, in the Z sources, the inner radius inferred from both the kHz QPOs and the accretion-disk spectral component have been determined to vary rapidly at a luminosity which is thought to be close to the Eddington value (\\cite{Yu,Lin}). In particular, the kHz QPOs vary with the normal branch oscillation (NBO) phase, which may indicate radiation induced rapid drifts of the innermost accretion flow. Perhaps these findings could be understood in terms of the expected behaviour of accreting matter accumulating on (a sector of) the Eddington sphere.  In the non-spherical geometry of accretion disks, radiation scattered in non-radial directions can escape from the system. Therefore the accreting fluid can ``levitate\" above the star only as long as it is optically thin. As more and more matter accumulates at the capture radius of Eq. (1), the shell eventually becomes optically thick and the outer layers of the fluid outweigh the radiation pressure support, leading to rapid accretion of matter and to the consequent evacuation of the region near the Eddington capture sphere, allowing the process of accumulation to begin again.  One would also expect the accreting fluid to be spread out to higher altitudes (lower and higher values of the polar angle $\\theta$) than those subtended by the inner parts of the optically thick disk, as the incoming fluid forms a ``puddle\" on the surface of the Eddington capture sphere.  A more detailed understanding of actual behaviour of Z sources would require relaxing the assumption of spherical symmetry, on which all the calculations in this paper were based. Comparison of the theoretical expectations with the observed NBO and its related flaring branch oscillation (FBO) would be very interesting.  In the atoll sources, which are less luminous, the Eddington luminosity is attained during X-ray bursts. It would be worthwile to investigate whether the Eddington capture sphere plays a role in the evolution of ``radius expansion bursts\" (e.g., \\cite{Damen}). In addition, in the persistent emission at rather high luminosity levels $\\sim 0.4\\, L_{\\rm Edd}$, a $\\sim 7$ Hz QPO similar to the NBO in the Z sources was also observed in the atoll source 4U 1820-30 (\\cite{Wijnands}). The effect we are considering above is global, but may apply to local properties of the accretion flow as well."
    },
    "1208/1208.3417_arXiv.txt": {
        "abstract": "Understanding the equation of state (EOS) of neutron-rich matter is a central goal of nuclear physics that cuts across a variety of disciplines.  Indeed, the limits of nuclear existence, the collision of energetic heavy ions, the structure of neutron stars, and the dynamics of core-collapse supernova all depend critically on the nuclear-matter EOS. In this contribution I focus on the EOS of cold baryonic matter with special emphasis on its impact on the structure, dynamics, and composition of neutron stars.  In particular, I discuss how laboratory experiments on neutron skins as well as on Pygmy and Giant resonances can help us elucidate the structure of these fascinating objects. ",
        "introduction": "\\label{introduction}  One of the four overarching questions framing the recent report by The Committee on the Assessment of and Outlook for Nuclear Physics is {\\sl ``How does subatomic matter organize itself?''}\\,\\cite{national2012Nuclear}.  This question has been at the core of nuclear physics since Rutherford's century-old discovery of the atomic nucleus in 1911.  The number of electrons---which equals the number of protons in a neutral atom---determines the chemistry of the atom. And it is this chemistry that is responsible for binding atoms into molecules and molecules into both traditional and fascinating new materials. But how does matter organize itself at densities significantly higher than those found in everyday materials; say, from $10^{4}\\!-\\!10^{15}\\,{\\rm g/cm^{3}}$.  Recall that in this units nuclear-matter saturation density equals $\\rho_{{}_{0}}\\!=\\!2.48\\times 10^{14}{\\rm g/cm^{3}}$. Indeed, relative to every day life these densities are so high that atoms become pressure ionized. Understanding what novels phases of matter emerge under these extreme conditions of density is both fascinating and unknown. Moreover, it represents one of the grand challenges in nuclear physics. Remarkably, most of these exotic phases---{\\sl Coulomb crystals, nuclear pasta, color superconductors}---can not be realized under normal laboratory conditions. Yet, whereas most of these phases have a fleeting existence in the laboratory, they attain stability in neutron stars due to the presence of enormous gravitational fields. In this manner neutron stars become the catalyst for the formation of unique states of matter and provide unique laboratories for the characterization of the ground state of cold matter over an enormous range of densities. Note that an unavoidable consequence of charge neutrality is that neutron-star matter is necessarily neutron rich. This is a natural consequence of the very low electron mass which in turn results in a high electron chemical potential.  \\begin{figure}[h] \\begin{center} \\includegraphics[height=3.5in]{Fig1.jpg} \\vspace{-0.2cm} \\caption{A scientifically-accurate rendition of the structure and the various phases predicted to exist in a neutron star (courtesy of Dany Page).} \\label{Fig1} \\end{center} \\end{figure} ",
        "conclusions": "\\label{conclusions}  Measurements of neutron radii provide important constraints on the isovector sector of nuclear density functionals and offer vital guidance in areas as diverse as atomic parity violation, heavy-ion collisions, and neutron-star structure. In this contribution we examined the possibility of using the quintessential nuclear mode---the isovector dipole resonance---as a promising complementary observable. For this mode of excitation in which protons oscillate coherently against neutrons, the symmetry energy acts as its restoring force. Thus, models with a soft symmetry energy predict large values for the symmetry energy at the densities of relevance to the excitation of this mode. As a consequence, softer models generates a dipole response that is both hardened and quenched relative to the stiffer models. However, being protected by the TRK sum rule, the energy weighted sum rule is largely insensitive to this behavior. In contrast, for the inverse energy-weighted sum---which is directly proportional to the electric dipole polarizability $\\alpha_{\\raisebox{-1pt}{\\tiny D}}$---the quenching and hardening act in tandem.  Thus, models with a soft symmetry energy predict smaller values of $\\alpha_{\\raisebox{-1pt}{\\tiny D}}$ than their stiffer counterparts. This results in a powerful ``data-to-data'' relation: {\\sl the smaller $\\alpha_{\\raisebox{-1pt}{\\tiny D}}$ the thinner the neutron skin}.  A particular intriguing question concerns the role of the pygmy dipole resonance in constraining the density dependence of the symmetry energy.  Regarded as an oscillation of the neutron-rich skin of a heavy nucleus against its isospin-symmetric core, the PDR was suggested to be strongly correlated to the neutron skin. In the particular case of the Tin isotopes, a clear emergence of low-energy dipole strength is observed as the nucleus develops a neutron-rich skin. Moreover, it appears that although the total EWSR is fairly insensitive to the density dependence of the symmetry energy, the fraction of the EWSR exhausted by the pygmy displays some sensitivity. However, in the case of the dipole polarizability the conclusion that the PDR is highly sensitive to the density dependence of symmetry energy appears inescapable. Indeed, in the particular case of ${}^{68}$Ni the PDR accounts for 20-25\\% of the total dipole polarizability and displays a strong sensitivity to the neutron skin. Yet, many open questions remain. First and foremost, the strong correlation between the PDR and the neutron skin found here appears to be model dependent. While we support the notion of a strong correlation between these two observables, Reinhard and Nazarewicz conclude that the neutron-skin thickness of ${}^{208}$Pb is very weakly correlated to the low-energy dipole strength\\,\\cite{Reinhard:2010wz}. Moreover, even the nature of the low-energy mode is unclear. Is it indeed a collective mode? Is it a skin oscillation? Can it be cleanly decoupled from the low-energy tail of the giant resonance? Although most of these issues were not addressed in this contribution, attempts to answer some of these question may be found in two recent reviews\\,\\cite{Paar:2007bk,Paar:2010ww}. Regardless of the nature of the mode, the emergence of low-energy dipole strength as nuclei develop a neutron-rich skin is an incontrovertible fact. As such, it should play a pivotal role in constraining the EOS of neutron-rich matter.  In summary, motivated by two seminal experiments\\,\\cite{Abrahamyan:2012gp,Tamii:2011pv}, we examined possible correlations between the electric dipole polarizability and the neutron skin of neutron-rich nuclei. The neutron-skin thickness of a heavy nucleus is a quantity of critical importance for our understanding of a variety of nuclear and astrophysical phenomena.  In particular, the neutron-skin thickness of $^{208}$Pb can provide stringent constrains on the density dependence of the symmetry energy which, in turn, has a strong impact on the structure, dynamics, and composition of neutron stars.  We conclude that precise measurements of neutron skins and $\\alpha_{\\raisebox{-1pt}{\\tiny D}}$---ideally on a variety of nuclei--- should significantly constrain the isovector sector of the nuclear energy density functional and will provide critical insights into the nature of neutron-rich matter."
    },
    "1208/1208.4830_arXiv.txt": {
        "abstract": "A new algorithm developed to perform autonomous fitting of gravitational microlensing lightcurves is presented. The new algorithm is conceptually simple, versatile and robust, and parallelises trivially; it combines features of extant evolutionary algorithms with some novel ones, and fares well on the problem of fitting binary-lens microlensing lightcurves, as well as on a number of other difficult optimisation problems.  Success rates in excess of $90\\%$ are achieved when fitting synthetic though noisy binary-lens lightcurves, allowing no more than 20 minutes per fit on a desktop computer; this success rate is shown to compare very favourably with that of both a conventional (iterated simplex) algorithm, and a more state-of-the-art, artificial neural network-based approach. As such, this work provides proof of concept for the use of an evolutionary algorithm as the basis for real-time, autonomous modelling of microlensing events. Further work is required to investigate how the algorithm will fare when faced with more complex and realistic microlensing modelling problems; it is, however, argued here that the use of parallel computing platforms, such as inexpensive graphics processing units, should allow fitting times to be constrained to under an hour, even when dealing with complicated microlensing models. In any event, it is hoped that this work might stimulate some interest in evolutionary algorithms, and that the algorithm described here might prove useful for solving microlensing and/or more general model-fitting problems. ",
        "introduction": "Gravitational microlensing is an established technique for detecting exoplanets. When a massive foreground object (the lens, e.g.\\ a planet and its host star) passes in front of a distant, background star (the source), the latter is magnified and displays a characteristic microlensing lightcurve. Since the lens is detected on account of its mass rather than its luminosity, faint planetary-mass objects can be detected by microlensing. Indeed, microlensing has the potential to yield the most representative statistical sample of Milky Way planets -- unlike many complementary techniques used to detect exoplanets, it is in principle sensitive enough to detect even very distant, Earth-mass planetary objects \\citep{Bennett:1996,Wambsganss:2011}. Unfortunately, microlensing events are extremely rare, requiring a very precise alignment between observer, lens and source: as of July 2012, of the several hundred known exoplanets, only around a dozen were discovered by microlensing (see \\citealp{Shvartzvald:2012}, and references contained therein). Still, many of these detections have constituted important discoveries in the broader context of exoplanetary science \\citep[e.g.][]{Beaulieu:2006,Gaudi:2008,Cassan:2012} -- \\textcolor{black}{recently, it has even been suggested that microlensing has facilitated the detection of a number of free-floating, planetary mass objects in the Galaxy (\\citealt{Sumi:2011}; cf., however, \\citealt{Quanz:2012}).}  Very comprehensive models do exist for describing microlensing events and their corresponding lightcurves, though unfortunately it is notoriously difficult to use these models to interpret microlensing events: amongst other complicating factors, the models tend to be highly nonlinear, and have enormous parameter spaces that are often fraught with ambiguities and degeneracies \\citep{Dominik:1999,Vermaak:2007}. Even the simplest possible microlensing model (viz.\\ a point-like source star lensing an isolated mass) poses some nontrivial challenges to microlensing modellers \\citep{Dominik:2008,Dominik:2009}.  This paper presents a new metaheuristics algorithm which combines features of extant evolutionary algorithms (genetic algorithms; evolution strategies) with some novel ones, developed with a view to performing efficient and autonomous fitting of (especially binary-lens) microlensing lightcurves.  The algorithm is, however, robust enough to solve general nonlinear optimisation problems, and its development was informed by tests carried out on a broad class of optimisation problems.  Section 2 of the paper gives an overview of the binary-lens model used to test the performance of the new algorithm; Section 3 describes evolutionary algorithms in general, as well as the new algorithm which is the focal point of this work (Appendix A, at the very end of the paper, provides a more detailed look at the `nuts and bolts' of the new algorithm); Section 4 focuses on the fitting experiments (and their results) used to assess the new algorithm; and Section 5 contains a commentary on the results presented in the preceding section. Section 6 concludes. \\section[Overview of binary-lens model]{Overview of binary-lens model} \\begin{figure} \\begin{center} \\includegraphics[scale=0.85]{Images/LensGeom.eps} \\caption{Geometry assumed in the SBLM.} \\label{fig:lensGeom} \\end{center} \\end{figure} A small companion to a stellar lens can be detected via perturbations it introduces to the lightcurve expected for an isolated lens. The resulting binary-lens lightcurve can exhibit a very wide variety of morphologies: it might be practically indistinguishable from a point-lens lightcurve, or it might exhibit complex structure including significant asymmetry, multiple peaks, and spikes of high (formally infinite) magnification produced during so-called `caustic crossings' \\citep*{Mao:1991,Night:2008}. Models of such binary-lens events are particularly useful for characterising exoplanetary microlensing events: even a system with one star and multiple planetary bodies can often be well-approximated either by ignoring multiple planets, or by treating each planet plus its host star as an independent binary system, provided source magnification is not too high \\citep*{Gaudi:1998}.  The model introduced here features $7$ basic parameters that describe rectilinear motion of an unblended, point-like source across a static binary lens. This simple binary-lens model (hereafter SBLM) neglects some higher-order effects that need to be taken into account when carrying out in-depth modelling of binary-lens events: nevertheless, the model is far from trivial, and is useful enough to provide first-order fits to many binary-lens events.  For example, even though the assumption of a point-like source will break down in a large fraction of real events \\citep{Dominik:2008}, the regions of lightcurves affected by finite-source effects tend to be localised -- usually affecting the high-magnification peaks associated with e.g.\\ caustic crossings -- so even when such effects are present, good SBLM-based fits can often be obtained simply by excluding the relevant regions of lightcurves from the fitting \\citep{Vermaak:2007}. It should be emphasised, however, that this work does not attempt to advocate the SBLM for use in modelling real microlensing events: the model is adopted here primarily to provide a relatively straightforward platform for benchmarking different algorithms studied later in this work (see Section \\ref{section:experiments}). \\subsection{Model parameters} A useful scale for characterising a lensing event is the so-called Einstein radius or Einstein angle of the primary lens, $\\theta_E$, defined as: \\begin{equation} {\\theta _E} := \\sqrt {\\frac{{4GM}} {{{c^2}}}\\left( {\\frac{{{D_S} - {D_L}}} {{{D_L}{D_S}}}} \\right)}, \\end{equation} where $G$ is the gravitational constant, $M$ is the mass of the primary lens, $c$ is the vacuum speed of light, $D_L$ is the distance between observer and the lens, and $D_S$ is the distance between observer and the source. The Einstein radius is the angular radius of the ring that would be formed in the case of perfect source-lens-observer alignment; for typical Bulge microlensing events, $\\theta_E\\sim1~\\textrm{mas}$.  For convenience, complex notation is used here to describe the lensing event. The origin of the coordinate axes is placed at the projected position of the primary lens; with no loss of generality the secondary lens is placed on the real ($+x$) axis; and $\\zeta\\in\\mathbb{C}$ is used to denote the position of the source on the sky.  The SBLM's seven parameters, and their physically-permissible ranges, then, are as follows (refer also to Fig.\\ \\ref{fig:lensGeom}): \\begin{enumerate} \\item $a$, the projected angular orbital separation, in units of $\\theta_E$, between the primary and secondary lenses ($0<a<\\infty$); \\item $b$, the length, in units of $\\theta_E$, of the projected source-lens impact vector $\\bmath{b}$ ($0<b<\\infty$); \\item $m_0$, the unlensed magnitude of the source; \\item $q$, the ratio of the mass of the secondary lens to that of the primary lens ($0<q<1$); \\item $\\theta$, the angle formed between the $+x$ axis and the projected impact vector $\\bmath{b}$ ($0<\\theta<2\\pi$); \\item $t_E$, the Einstein radius crossing time, i.e.\\ the \\textcolor{black}{time it takes} the source to travel an angular distance of $\\theta_E$; and \\item $t_m$, the time of closest projected approach of the source to the primary lens. \\end{enumerate} Note that to obtain physical angular distances, the parameters $a$ and $b$ must be scaled by $\\theta_E$.  Many other parametrizations of binary lensing events are possible \\citep[see e.g.][]{Mao:1995,Dominik:1999}, although most are qualitatively similar to the one used here. In all subsequent discussions, the seven SBLM parameters will be assumed to be constrained to the ranges given in Table \\ref{tab:params}. These ranges were chosen to facilitate direct comparison with the results of \\citet{Vermaak:2003,Vermaak:2007}; they cover most of the lensing zone, and thus the geometries with the highest likelihood of detecting secondary lenses.   \\begin{table} \\centering \\caption{Allowed parameter ranges for all fits and simulated lightcurves.} \\label{tab:params} \\begin{tabular}{cccc} \\hline Parameter & Units &    Minimum &    Maximum \\\\ \\hline $a$ & $-$ &      $0.6$ &      $1.7$ \\\\  $b$ & $-$ &    $0.001$ &        $1$ \\\\  $m_0$ & mag &       $18$ &       $22$ \\\\  $q$ & -- &      $0.1$ &        $1$ \\\\  $\\theta$ & rad &        $0$ &      $2\\pi$ \\\\  $t_E$ & d &        $5$ &       $50$ \\\\  $t_m$ & d &      $-20$ &       $20$ \\\\ \\hline \\end{tabular} \\end{table}  \\subsection{Calculation of binary-lens lightcurves} One of the primary difficulties associated with the SBLM is that it does not provide an explicit expression for the amplification of the source as a function of time and lensing parameters. Instead, the time-varying source position relative to the lens, as well as the lens geometry, is first used to calculate the position of the \\emph{images} of the source that are formed by the lens; the light from each of these images is then summed to calculate the total amplification.  Assuming rectilinear motion, straightforward geometric considerations yield the \\emph{source} position as a function of time, and two of the lensing parameters: \\begin{equation}\\label{eq:sourcePos} \\zeta  = \\tau\\sin \\theta  + b\\cos \\theta  + i\\left( { - \\tau\\cos \\theta  + b\\sin \\theta } \\right), \\end{equation} where $\\tau$ is a dimensionless time parameter: \\begin{equation} \\tau := \\frac{{t - {t_m}}}{{{t_E}}}. \\end{equation} Assuming zero external gravitational shear and a binary lens, the general lens equation \\citep*[see][]{Bourassa:1973,Witt:1990} reduces to \\begin{equation}\\label{eq:lensing} \\zeta  = z - \\frac{1}{{\\bar z}} + \\frac{q}{{a - \\bar z}}, \\end{equation} which, with $\\zeta$ known, is an implicit expression for the \\emph{image} positions, $z\\in\\mathbb{C}$. It may be shown that this special case of the lensing equation has always either $n=3$ or $n=5$ solutions, corresponding to 3 or 5 lensed images. Following some nontrivial algebraic manipulation \\citep{Witt:1990}, Eqn.\\ \\ref{eq:lensing} can be converted into a complex quintic polynomial equation of the form \\begin{equation}\\label{eq:lensingPoly} f(z;\\zeta ,a,q) = \\sum\\nolimits_{j = 0}^5 {{c_j}(\\zeta ,a,q)  {z^j} = 0}, \\end{equation} where $c_j\\in\\mathbb{C}$; the roots of this polynomial can be found using any one of a number of standard numerical techniques, for example the Jenkins-Traub algorithm or Laguerre's algorithm (in practice it is more convenient to solve this polynomial equation than trying directly to find all the solutions to Eqn.\\ \\ref{eq:lensing}). Two of the five polynomial roots may not be true solutions to the lensing equation itself, but these spurious solutions may readily be eliminated by direct substitution into Eqn.\\ \\ref{eq:lensing}.  In an alternative though equivalent formulation, the lensing equation can be cast into the form of a real polynomial equation \\citep{Asada:2002}.  Once the image positions $z_i$ have been solved for, the amplification due to the $i^{\\rmn{th}}$ lensed image is calculated via \\begin{equation} {A_i} = \\frac{1}{{\\det \\left| {{\\mathsf{J}_i}} \\right|}} = {\\Biggl| {1 - \\biggl| {\\frac{1}{{{z_i}^2}} + \\frac{q}{{{{(a - {z_i})}^2}}}} \\biggr|} \\Biggr|^{ - 1}}, \\end{equation} where $\\mathsf{J}_i$ is the Jacobian of Eqn.\\ \\ref{eq:lensing} for the case of the $i^{\\rmn{th}}$ image. Finally the amplification due to the $n$ individual images is summed, so that the lensed magnitude $m$ is given by: \\begin{equation}\\label{eq:A_s} m^{(k)} =  - 2.5\\cdot{\\log _{10}}\\left( {\\sum\\nolimits_{i = 1}^n {{A_i^{(k)}}} } \\right) + {m_0}; \\end{equation} the superscript $k=1,2,\\ldots,N$ is introduced to emphasise that the preceding calculations culminate in a single point on an $N$-point lightcurve.  \\subsection{Difficulties associated with fitting binary lightcurves}\\label{sec:difficulties} Given a set of parameters, obtaining the lightcurve predicted by the SBLM is a nontrivial though relatively straightforward task. Unfortunately, however, the associated \\emph{inverse problem} -- i.e., given a lightcurve, using the SBLM to determine the physical parameters associated with the lensing event -- is extremely difficult; yet it is of course the solution to this inverse problem that is of value insofar as the interpretation of real-world microlensing data is concerned. Some of the difficulties associated with this fitting problem include the following. \\begin{enumerate} \\item \\emph{Computational complexity}. Generating even a single $N$-point lightcurve, i.e.\\ solving the \\emph{forward problem}, requires, as a minimum, the roots of $N$ different quintic polynomials to be found numerically \\citep[for techniques on solving quintic polynomial equations, see e.g.][]{King:1996,Ralston:2001,Press:2007}. \\item \\emph{Volume of parameter space}. The SBLM has a 7-dimensional parameter space, and the physically-realistic ranges for some of the parameters (e.g.\\ impact parameter, lens mass ratio) can span many orders of magnitude.  \\citet{Vermaak:2003} estimated that an exhaustive grid search of the parameter space would take on the order several of years to complete (assuming a $2\\%$ error tolerance on all parameters, the very conservative parameter ranges in Table \\ref{tab:params}, and $\\sim2$~ms/lightcurve calculation). Any extensions to the SBLM only exacerbate this problem. \\item \\emph{Nonlinearity}. The mapping from SBLM parameter space to source amplification is highly nonlinear -- consequently, nearly-identical parameter sets can give rise to dramatically different lightcurves. A typical regression surface will contain a large number of local optima, will be non-smooth, and will contain no clue as to where the global optima (with respect to some, e.g.\\ $\\chi^2$-based, metric) are to be found; furthermore the wells of convergence around optima tend to be small \\citep{Vermaak:2003,Bennett:2010}, and when dealing with noisy data, true solutions need not even correspond to a globally optimal solutions. Using biased parameter estimators (e.g.\\ using a maximum likelihood estimator, rather than, say, a posterior-mode estimator) can also lead to globally optimal solutions that are, potentially, nowhere near the true solution  (\\citealp{Dominik:2008}; see also Section \\ref{section:experiments} of this paper). \\item \\emph{Degeneracy}. The aforesaid problem of similar parameters leading to very different lightcurves can be overcome by making a denser sampling of the parameter space; more challenging, however, is the problem that a number of very \\emph{different} parameter sets can give rise to virtually indistinguishable lightcurves \\citep{Dominik:1999}. In practice, such degeneracies can either be resolved by using non-photometric (e.g.\\ spectroscopic) data to reduce the effective parameter space, or they can be avoided with high quality photometry and dense lightcurve sampling \\citep{Mao:1995,Bennett:2008}.\\footnote{Some of the degeneracies identified by \\citet{Dominik:1999} will not even be resolved with any currently-achievable data qualities and sampling rates \\citep{Shvartzvald:2012}.} \\end{enumerate} The impasse, then, is that there is little hope of finding the `correct' set of parameters to describe a microlensing event if one does not make a dense and extensive search of the parameter space, but the computational burden of doing so is often crippling. A thorough  search is especially important because even if one does happen to find one set of parameters that seems to provide a very good description of the event in question, there may exist many other parameter sets that provide equally good or better descriptions of the event. A good search algorithm should, therefore, be able to locate, in a reasonable amount of time, all \\emph{potential} solutions; actually choosing between competing model parameter sets is a different problem altogether \\citep{Burnham:2002}. \\begin{figure*} \\begin{center} \\includegraphics[scale=0.75]{Images/EA} \\caption{Schematic to illustrate the workings of the canonical genetic algorithm, where trial solutions are encoded as binary strings. Each bit represents a gene; here, lighter shades are used to represent solutions determined to be fitter, according to any problem-specific metric. The `initial' population (first array) contains a range of solutions, some good (lighter shades) and some bad (darker shades). During ranking, the solutions are sorted from best to worst (second array). The genes from fitter solutions are then favoured during recombination, which is apparent from the fact that the new solutions (third array) are assembled mainly from components of the original lighter-coloured solutions; finally, random mutations -- which may or may not be beneficial -- are introduced (fourth array), and the whole process repeats.  The fine-grained details of many EAs differ from the simple algorithm illustrated here; however, most EAs share the same underlying concept of a population of trial solutions evolving, under the action of a few simple evolutionary operators, towards optimality (the constitution of the or an `optimal' solution will depend on the problem at hand).} \\label{fig:EA} \\end{center} \\end{figure*} \\subsection{A simple noise model} A simple model of photometric noise was developed, during the course of this work, based on fits to $1000$ microlensing events observed during the 2011 campaign of the Optical Gravitational Lensing Experiment \\citep{Udalski:2003}; the full dataset comprised approximately \\hbox{$1.5$~million} \\hbox{$I$-band magnitudes}, along with estimated photometric errors, seeing estimations and sky levels.\\footnote{The data are available online at http://ogle.astrouw.edu.pl.} The best-fitting model for photometric errors as a function of $I$-band magnitude (based on data in the range \\hbox{$14\\la m_I^{(k)}\\la22$}) thus obtained was \\begin{equation} \\label{eq:noiseModel} {\\log _{10}}{\\sigma _I^{(k)}} = 0.3416 \\cdot {m_I^{(k)}} - 7.7095, \\end{equation} where $\\sigma _I^{(k)}$ is to be interpreted as the standard deviation of a Gaussian distribution of observed magnitudes around some unobserved true magnitude $m_I^{(k)}$.  This model seems overly optimistic for bright sources (i.e.\\ $m_I^{(k)}\\sim14$, where the model predicts photometric errors on the order of $0.1\\%$!), and even the assumption of Gaussian noise is probably na\\\"ive \\citep{Dominik:2008}; still, the model does at least provide a means of making the simulated lightcurves used for fitting experiments (Section \\ref{section:experiments}) somewhat more realistic and more challenging to fit. % \\section[Evolutionary algorithms]{Evolutionary algorithms} \\subsection{Overview} Evolutionary algorithms (hereafter EAs), are metaheuristic optimisation algorithms which tend to yield good results on a wide range of (even extremely difficult) mathematical optimisation problems. EAs draw inspiration from evolutionary biology -- especially population genetics -- and they incorporate, in a computational setting, notions such as natural selection/survival of the fittest, genetic recombination, inheritance, and mutation. The first EA-based mathematical optimiser was proposed in the mid-1970s \\citep{Holland:1975}, and since then, many modifications and improvements to the basic algorithm have been developed, including mechanisms without any direct biological analogues \\citep{Haupt:2004}.  So-called \\emph{genetic algorithms} (GAs) form one of the most successful subsets, and certainly the best-known subset, of evolutionary algorithms; \\emph{evolution strategies} (ESs), developed independently from (though more or less concurrently with) GAs, form another well-known subset. In spite of the rich variety of their potential incarnations, most EAs share a basic working scheme.  As an example, consider a typical GA. The algorithm starts with a large, randomly-generated population of candidate solutions (called individuals or phenotypes) to the optimisation problem at hand, and associates with each solution an encoded version of the phenotype (called a chromosome, genotype or an individual's genetic material), as well as a problem-specific measure of the solution's quality (fitness). Then, by repeated application of `genetic operators' mainly at the genotypic level, the algorithm causes the population as a whole to increase in phenotypic fitness -- that is, solutions evolve towards optimality.  A typical (though simplistic and by no means general or optimal) working scheme for a GA is as follows: \\begin{enumerate} \\item construct a random initial population of genotypes; \\item decode the genotypes and evaluate their phenotypic fitness; if the fittest phenotype matches the user-defined target fitness (or other termination criterion), terminate, otherwise continue; \\item produce offspring by stochastic selection and recombination of genetic material in the current population, favouring the genes of individuals with high phenotypic fitness; \\item introduce, with some low probability, random changes (copying errors) into the genetic material of the offspring; \\item replace low-fitness members of the old population with the offspring created in the previous step, and return to step (ii). \\end{enumerate} This scheme is illustrated in Fig.\\ \\ref{fig:EA}. The selective recombination of the population's genetic material exploits information associated with good solutions to build even better ones, and the random mutations serve to inject entirely new and potentially favourable material into the gene pool that could not be obtained simply by recombining the genetic material of existing individuals. Such an evolutionary scheme has a number of features which distinguish it from random heuristics, albeit that it might bear superficial resemblance to e.g.\\ a standard Monte Carlo approach \\citep{Gregory:2005}.  It may be shown that, subject to a few reasonable assumptions, the canonical GA will always converge to the global optima in the search space in question \\citep*{Eiben:1990,Michalewicz:1996}; moreover, it is also possible, though usually not straightforward, to estimate the \\emph{rates} at which solutions are likely to be located on given problems \\citep[e.g.][]{Holland:1975,Thierens:1994}.  ESs are similar to GAs in many respects, although usually they avoid solution encoding or discretisation, evolve the algorithm's control parameters in tandem with the trial solutions, and/or use more sophisticated mutation schemes \\citep{Kramer:2010}.  It should be noted that in many problems -- microlensing modelling included -- one usually cannot hope to arrive at a unique or clear-cut, globally-optimal solution, and one must instead try to sample all possible optima in a multimodal parameter space. Even if an EA is not specifically set up to search for multiple solutions, mutation operators ensure exploration of the entire parameter space (provided the evolutionary sequence is sufficiently long) -- as such, given a completed evolutionary sequence, multiple possible solutions can be obtained simply by extracting all explored phenotypes that meet some fitness criterion (as opposed to extracting only the fittest phenotype in the final population state, which should correspond to a global optimum). Moreover, if one knows beforehand that one is dealing with multimodal search spaces, simple modifications can be made to make EAs more efficient at searching for multiple solutions. A simple approach is to begin a new evolutionary sequence whenever the population starts to stagnate around an optimum; alternatively, multiple independent populations can be evolved concurrently. A more sophisticated approach might be to dedicate a fraction of the evolutionary population(s) to global exploration of the search space, and the remaining fraction purely to the exploitation of promising solutions already discovered \\citep{Tsutsui:1997}. \\subsection{Advantages of evolutionary algorithms} EAs in general offer a number of benefits that suggest their suitability for dealing with the challenges discussed in Section \\ref{sec:difficulties}; some of these are outlined below. \\begin{enumerate} \\item \\emph{Robustness}. EA-based optimisers can handle problems where the spaces to be searched for optima are multimodal, have low-contrast or have a very high dimensionality \\citep{Charbonneau:1995}.  \\item \\emph{Simplicity}. In order to solve a given optimisation problem, most `off-the-shelf' EAs require only a single, unambiguous measure of the quality (fitness) of candidate solutions. They do not require, for example, gradients or Hessian matrices, the computation of which might be prohibitively difficult or impossible in some problems.  \\item \\emph{Speed}. Apart from the intrinsically high speed with which EAs tend to explore large parameter spaces, they are embarrassingly parallel: that is, very little effort is required to transform a serial implementation into a parallel implementation. Thus they are well-suited to exploiting high-performance hardware such as multi-core workstations, graphics processing units (GPUs), clusters, etc.: see Section \\ref{sec:GPUs}. To be sure, canonical EAs (GAs, ESs, etc.) do have a reputation for being inefficient local optimisers \\citep{Charbonneau:1995,Michalewicz:1996}, but this problem can readily be mitigated either by coupling an EA to a dedicated local optimiser -- e.g.\\ a gradient-based method -- or, as suggested in Section \\ref{section:EMMA} below, by making simple tweaks to an EA's underlying mechanisms, so as to endow it with local optimisation capabilities.  \\item \\emph{Versatility}. A single EA-based optimiser can be expected to yield `good enough' results on a very wide class of problems -- from a problem as simple as fitting a three-parameter Gaussian to some data, to one as complex as choosing a molecular configuration to minimize a Buckingham potential with hundreds of parameters \\citep{Wehrens:1998} -- and it is easy to incorporate problem-specific knowledge into an EA-based solver \\citep{Haupt:2004}. \\end{enumerate}  The widespread adoption of EAs in fields such as engineering, chemistry, biology, and economics bears testimony to their many merits, and though their uptake in the physical sciences has been somewhat slower (at least partly because the theoretical understanding of their workings is still quite limited), recently they have already been used with great success in many branches of astronomy and astrophysics \\citep{Rajpaul:2011}. In fact, a classical genetic algorithm -- PIKAIA -- has already been tested in the context of microlensing modelling \\citep{Kubas:2005}, although it has not seen much use in recent years (at least, not in the microlensing community); unlike PIKAIA, though, the new algorithm presented here is tailored for high performance numerical optimisation, rather than for more general and pedagogical purposes. \\subsection{A new algorithm for fitting microlensing events}\\label{section:EMMA} The algorithm developed by the author for fitting microlensing lightcurves combines features of extant evolutionary algorithms (GSs, ESs) with some novel ones. The algorithm's main features are the following (a more detailed examination of the `nuts and bolts' of the algorithm is presented in Appendix A): \\begin{enumerate} \\item a real-valued representation of solutions is used (i.e.\\ parameter domains are not discretised, as would've been the case with e.g.\\ a binary-coded GA); \\item fitness-ranking can, in general, be performed according to arbitrary Bayesian priors and likelihood functions; \\item individuals are chosen for reproduction via a variant of the well-known \\emph{tournament selection} mechanism \\citep[e.g.][]{Miller:1995}; \\item reproductive pairings comprise exactly two parent solutions, and give rise to exactly two offspring; \\item the offspring are constructed in one of three ways -- as random interpolants of their parents, as exact duplicates of their parents, or as solutions where each parameter is drawn randomly but without modification from one of the two parents; \\item both coarse-grained (`jump') and fine-grained (`creep') mutation operators are used, facilitating respectively global and local optimisation, with the mutation rates being dynamically adjusted in response to changes in population fitness; and \\item the evolutionary sequence is restarted whenever the population stagnates around an optimum. \\end{enumerate}  The algorithm runs autonomously in the sense that, given a set of data and a general model (e.g.\\ the SBLM) to be fit to the data, the algorithm will search intelligently\\footnote{The algorithm's `intelligence' stems from its ability to exploit (via e.g.\\ the reproduction and fine-grained mutation operators) structure  in the search space to guide its search, as well its ability to adapt its own control parameters in response to the progress of the search.} through the model parameter space to find solutions compatible with the data, without requiring any intervention from or assistance on the part of the user.  The algorithm's development was informed by tests carried out not only on single- and binary-lens fitting problems, but also on a range of more general nonlinear optimisation problems \\citep{More:1981, Haupt:2004}. The use of the dual mutation operators and dynamically-adjusted mutation rates means that, unlike more traditional EAs, the algorithm fares well at both global exploration and local optimisation. The jump operator allows for large jumps through the entire parameter space, whereas the creep operator is more conservative, and introduces only small (log-uniformly distributed) perturbations to existing solutions. The hybridised operators used for offspring construction are also more compatible with local fine-tuning than the canonical genetic crossover operator, which is well known to be antithetical to local optimisation \\citep{Eiben:1998}. Finally, the algorithm was designed in an array programming paradigm (wherein all operators are applied to a single population matrix rather than to individual solutions) and, as such, the main algorithm is able to comprise fewer than $100$ lines of code when implemented in high-level language like \\textsc{Matlab} or Python. ",
        "conclusions": "A novel evolutionary algorithm was introduced and was shown to be well suited to fitting binary-lens microlensing lightcurves. Despite the difficulty of this fitting problem, the algorithm yielded excellent fitting accuracy whilst maintaining a relatively modest computational footprint, and offering a number of other desirable properties (robustness, versatility, trivial parallelisation, etc.). As such, this work provided proof of concept for the use of an evolutionary algorithm as the basis for real-time, autonomous modelling of microlensing events. It was noted, however, that further work is required to investigate how the algorithm will fare when faced with more complex and realistic microlensing modelling problems.  Indeed, this work could be extended in a number of ways, in particular by relaxing some of the assumptions made here, and also by investigating how an evolutionary algorithm will fare at trying to predict features in the lightcurves of \\emph{ongoing} microlensing events. Early investigations in this regard have been promising.  Though evolutionary algorithms can by no means be the `final word' in microlensing modelling, a simple evolutionary algorithm such as the one presented here could serve as a useful addition to the toolbox of a microlensing modeller, or indeed, of any astronomer working with difficult model-fitting problems."
    },
    "1208/1208.1762_arXiv.txt": {
        "abstract": "Understanding the interaction between galaxies and their surroundings is central to building a coherent picture of galaxy evolution. Here we use {\\em GALEX} imaging of a statistically representative sample of 23 galaxy groups at $z\\approx 0.06$ to explore how local and global group environment affect the UV properties and dust-corrected star formation rates of their member galaxies. The data provide star formation rates out to beyond $2R_{200}$ in all groups, down to a completeness limit and limiting galaxy stellar mass of $0.06$~$M_\\odot$~yr$^{-1}$ and $1\\times 10^8$~$M_\\odot$, respectively. At fixed galaxy stellar mass, we find that the fraction of star-forming group members is suppressed relative to the field out to an average radius of $R\\approx 1.5$~Mpc\\,$\\approx 2R_{200}$, mirroring results for massive clusters. For the first time we also report a similar suppression of the specific star formation rate within such galaxies, on average by 40\\% relative to the field, thus directly revealing the impact of the group environment in quenching star formation within infalling galaxies. At fixed galaxy density and stellar mass, this suppression is stronger in more massive groups, implying that both local and global group environment play a role in quenching. The results favor an average quenching timescale of $\\ga 2$~Gyr and strongly suggest that a combination of tidal interactions and starvation is responsible. Despite their past and ongoing quenching, galaxy groups with more than four members still account for at least $\\sim 25$\\% of the total UV output in the nearby universe. ",
        "introduction": "Galaxies may evolve from star-forming late-types to passive early-types either via internal, secular processes or through mechanisms induced by their environment. Disentangling these two pathways and understanding the nature and impact of various environmental processes on galaxy evolution is a major goal of contemporary astrophysics. While the overall star formation density of the universe has declined significantly since a redshift of $z\\approx 2$ \\citep{mada96,hopk06}, this evolution appears to be accelerated in dense environments, with groups and clusters of galaxies containing a lower fraction of star-forming galaxies at fixed stellar mass and redshift than the general field (e.g., \\citealt{kauf04,mcge11}).  Recent work has demonstrated that the effect of secular evolution in quenching star formation, primarily driven by galaxy stellar mass, can be separated from that of the galaxy environment \\citep{peng10}. The former effect dominates at high galaxy masses, whereas environmental quenching of star formation becomes increasingly important at lower masses. Possible mechanisms for a more rapid quenching of low-mass galaxies in dense environments include tidal interactions, harassment, ram pressure stripping, starvation, and major and minor mergers.  Groups of galaxies are particularly interesting in this respect. Containing around half of all galaxies and most of the stellar mass in the local universe \\citep{eke04,eke05}, they represent a key environment in the hierarchical build-up of cosmic structure. The galaxy population within groups has evolved significantly since $z\\approx 0.5$, with the fraction of emission-line galaxies having declined substantially \\citep{wilm05}. Interestingly, however, the emerging picture is one in which star formation in star-forming galaxies of a given stellar mass is similar in groups and the field, whereas the fraction of such galaxies is not (e.g., \\citealt{balo04,bald06,iovi10}). For example, \\citet{balo04} used inferred H$\\alpha$ equivalent widths of galaxies within a large sample of low-redshift groups in the Sloan Digital Sky Survey (SDSS) and 2dF Galaxy Redshift Survey (2dFGRS) catalogs to show that the fraction of star-forming galaxies depends systematically on local density, but the actual star formation rate {\\em in} such galaxies does not. This has recently been shown to apply also at intermediate redshifts, $z \\sim 0.4$ and at fixed stellar mass \\citep{mcge11}, and the result has further been extended to galaxies in both nearby \\citep{hain11} and distant ($z \\sim 1$; \\citealt{muzz12}) clusters and superclusters.  This has been interpreted as evidence that the truncation of star formation is accomplished on short timescales, rapidly enough to leave the average star formation properties of star-forming galaxies largely unaffected. This would be true even in groups of fairly low mass, and support for this comes from recent studies based on UV star formation rates in Hickson Compact Groups which are generally relatively sparse systems. Such work has revealed a pronounced deficit of galaxies with moderate specific (i.e., stellar-mass-normalized) star formation rates \\citep{tzan10,walk12}, consistent with a rapid, environment--driven transition from (perhaps temporarily enhanced) star formation to quiescence. Although the dominant responsible mechanism(s) have yet to be unambiguously identified, it is clear that galaxy--galaxy interactions could play a prominent role both in compact and more typical groups, given the relatively low velocity dispersions characteristic of these environments.  Recognizing that a detailed multi-wavelength characterization of the coupling between global group properties and those of the group members themselves could offer important new insights, we initiated the {\\em XI} Groups Survey, a study of a statistically representative sample of 25 redshift-selected groups at $z\\approx 0.06$ \\citep{rasm06b}. In this framework, we have previously presented star formation rates (SFRs) from {\\em Spitzer}/MIPS 24\\,$\\mu$m data for a subset of nine groups \\citep{bai10}. Results revealed star-forming galaxy fractions that bridge those of the field and massive clusters, being everywhere higher than those in cluster outskirts but lower than in the field by $\\sim 30$\\% on average. The fractions showed no systematic dependence on global group properties such as velocity dispersion or total stellar mass. In addition, and in line with the above previous studies, specific SFRs were generally found to be similar to those in the field.  One limitation of our {\\em Spitzer} study was the narrowness of the rectangular region ($20\\arcmin$, equivalent to $\\sim 0.7$~Mpc radius at the sample redshift) covered by our MIPS observations. To fully understand the suppression of star formation in groups requires establishing the star-forming properties of group members out to larger radii, and including galaxies that are encountering the group environment for the first time. For example, simulations \\citep{kawa08} and isolated observations \\citep{rasm06a} suggest that gaseous stripping and starvation can suppress star formation in disk galaxies on their first passage through even low-mass groups. In more massive clusters, star formation rates appear suppressed out to $R\\sim 2R_{200}$ (\\citealt{balo98}; here $R_{200}$ is the radius enclosing a mean overdensity of 200 relative to the critical density). Recent studies have extended this conclusion to even larger radii \\citep{chun11}, revealing a suppression in the fraction of star-forming galaxies out to $\\sim 7$~Mpc from cluster cores \\citep{lu12}. However, other works have also shown that star formation can be locally enhanced well outside the cluster virial radius for galaxies within the filaments feeding these structures \\citep{port07,fadd08,port08,pere10}, with galaxy--galaxy interactions among infalling galaxies providing a possible explanation. Testing the general validity of these results for low-mass groups requires full coverage of the group environment out to the field and infall regions.  Here we extend our previous {\\em Spitzer} study by presenting {\\em GALEX} data for most of the full {\\em XI} sample. Apart from providing improved statistics due to the increased sample size, the {\\em GALEX} field-of-view of $R=36\\arcmin$ also allows coverage of each group field out to $R\\sim 2.5$~Mpc, corresponding to at least $2R_{200}$ for these systems. This continuous coverage of all regions from the general \"field\" through the infall regions to the dense group cores enables a complete census of the star-forming properties of group galaxies, including those only now encountering the group environment.  We assume $H_0=72$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_m=0.27$, and $\\Omega_{\\Lambda}=0.73$. At the sample median redshift of $z=0.061$, $1\\arcmin$ corresponds to 69~kpc and the total {\\em GALEX} field to $\\approx 5.0$~Mpc diameter. All uncertainties are quoted at the $1\\sigma$ level. ",
        "conclusions": "\\label{sec,summary}  Using {\\em GALEX} imaging, we have presented dust-corrected star formation rates and general UV properties of the galaxy population within 23 redshift-selected groups at $z\\simeq 0.06$. The data cover a radius of 2.5~Mpc around each group, allowing us to establish UV properties of group members from group cores into the general field, down to an SFR completeness level of 0.06~$M_\\odot$~yr$^{-1}$ before dust corrections. Our conclusions may be summarized as follows.  The fractions of UV star-forming galaxies show a clear, systematic decline toward the dense group cores. The fractions are suppressed by 35\\% on average relative to the field, and the suppression is significant out to at least $R \\approx 1.5$~Mpc~$\\approx 2R_{200}$. This implies that the group environment is modifying galaxy properties out to similar overdensity radii as are much more massive clusters \\citep{balo98}. This remains true at fixed galaxy stellar mass but is most pronounced for galaxies within the low-mass half of our UV--detected sample ($M_\\ast < 10^{9.8}$\\,$M_\\odot$). The star-forming fractions within $R_{200}$ also decline with total group stellar mass, from an average of $\\sim 70$\\% in the lowest-mass groups to only $\\sim 30$\\% within the most massive systems ($\\sigma \\ga 400$~km~s$^{-1}$).  In a minority of the groups, the central galaxy is a late-type object with NUV--R colors consistent with ongoing star formation. These groups tend to be of lower total stellar mass and richness, and are possibly dynamically younger than the rest, as also suggested by the larger average projected distance between their central galaxy and the group luminosity centroid.  The star formation rates of group galaxies depend primarily on galaxy stellar mass, in a manner broadly agreeing with results for the global low-redshift galaxy population, SFR\\,$\\propto M_\\ast^{\\sim 0.6}$. There is no evidence for enhanced star formation at intermediate radii or galaxy densities, as otherwise reported for galaxies in the filaments connecting massive clusters. At stellar masses $M_\\ast > 10^{8.5}$\\,$M_\\odot$, the global fraction of starburst galaxies within the groups is $\\approx 3$\\%, consistent with, but at the low end of, the range inferred globally for the nearby universe. The fraction declines significantly with local galaxy density, however, from 5--7\\% in the lowest-density regions to $\\la 1$\\% in the dense group cores.  For the first time, we report a dependence on local galaxy environment of not only star-forming galaxy fractions but also of specific SFRs {\\em within} blue star-forming galaxies at fixed $M_\\ast$. The latter trend is strongest for low-mass galaxies ($M_\\ast \\la 10^9$\\,$M_\\odot$), in which the mean specific SFR declines by a factor $\\approx 2.5$ from the field to the dense group cores. This effect becomes insignificant at galaxy stellar masses $M_\\ast \\ga 10^{10}$\\,$M_\\odot$, plausibly explaining why it has not been detected in other detailed studies of groups at higher redshift. On average, specific star formation rates of star-forming galaxies in groups are suppressed by $\\approx 40$\\% relative to the field, taking results for the field either from our own analysis or from other {\\em GALEX} results for large, low-redshift galaxy samples.   At fixed galaxy mass and local density, the suppression of both specific SFRs and star-forming galaxy fractions is stronger in more massive groups, especially in the dense group cores. The decline in both quantities with local density is in fact largely driven by the high-mass groups with red central galaxies, whereas most galaxies in lower-mass groups with blue centrals may have yet to be significantly affected by their environment. Both local and global group environment thus play a role in quenching star formation, suggesting a combination of at least two underlying mechanisms: One whose efficiency depends primarily on radius or local galaxy density, and one which scales with total group mass and acts primarily in the group cores. Galaxy--galaxy interactions and gas-dynamical processes, respectively, are the obvious candidates.  The observed decline in specific SFRs toward group cores implies a characteristic timescale for shutting down star formation which is comparable to the crossing times, $\\sim 2$~Gyr. This timescale is in good agreement with other independent estimates. Based on the observed starburst fraction in the groups, we demonstrate that gas consumption associated with starbursts triggered by galaxy--galaxy interactions proceed too quickly to be compatible with this timescale. Combined with the absence of a bimodality in the distribution of specific SFRs, this result confirms that a more slowly acting mechanism also contributes significantly to quenching in groups. We argue that starvation is the most likely candidate, acting in concert with galaxy--galaxy interactions.  The generality of these results for the global low-redshift galaxy population is attested by the finding that the total UV luminosity density provided by typical groups with more than four members is at least 25\\% in the local universe. This is in spite of the accelerated quenching of star formation relative to the field still taking place within these systems at present. We also find that $\\sim 70$\\% of the UV star formation in such group environments is obscured, in good agreement with results obtained for field and cluster samples."
    },
    "1208/1208.3551_arXiv.txt": {
        "abstract": "We study the distinction and quantification of chaotic and regular motion in a time-dependent Hamiltonian barred galaxy model. Recently, a strong correlation was found between the strength of the bar and the presence of chaotic motion in this system, as models with relatively  strong bars were shown to exhibit stronger chaotic behavior compared to those having a weaker bar component. Here, we attempt to further explore this connection by studying the interplay between chaotic and regular behavior of star orbits when the parameters of the model evolve in time. This happens for example when one introduces linear time dependence in the mass parameters of the model to mimic, in some general sense, the effect of self-consistent interactions of the actual N-body problem. We thus observe, in this simple time-dependent model also, that the increase of the bar's mass leads to an increase of the system's chaoticity. We propose a new way of using the Generalized Alignment Index (GALI) method as a reliable criterion to estimate the relative fraction of chaotic vs.~regular orbits in such time-dependent potentials, which proves to be much more efficient than the computation of Lyapunov exponents. In particular, GALI is able to capture subtle changes in the nature of an orbit (or ensemble of orbits) even for relatively small time intervals, which makes it ideal for detecting dynamical transitions in time-dependent systems. ",
        "introduction": "\\label{Intro}  The study of chaotic and regular properties of the motion in Hamiltonian systems constitutes a vast area of research in the field of nonlinear dynamics. Since the early 1960's, several methods and tools for the fast and accurate detection of the nature of orbits have been proposed and applied to this end in a great number of publications. One may refer e.g.~to the pioneering paper by H\\'{e}non and Heiles \\cite{HenHei:1964}, where the Poincar\\'{e} Surface of Section (PSS) \\cite[section 1.2b]{LL} was used to reveal the chaotic properties of a non-integrable 2 degree of freedom (2 dof) Hamiltonian system. Of great importance in this direction was also the algorithm proposed by Benettin and co-workers \\cite{Benettin1980a,Benettin1980b,SkoLE:2010} regarding the computation of the full spectrum of Lyapunov exponents (LEs) associated with the time evolution of deviation vectors from a reference orbit, which applies to dynamical systems of arbitrary dimension. More recently, other related methods have been proposed in the literature, like the ``Fast Lyapunov Indicator\" \\cite{FroLeGo:1997,FroLe:1998} and the ``Mean Exponential Growth of Nearby Orbits\" (MEGNO) \\cite{CinSimo:2000,CinGioSimo:2003}, while there have also been approaches focusing on the time series constructed by the coordinates of each orbit, like the ``Frequency Map Analysis\" \\cite{Las:1990,LasFroCel:1992,Las:1993} and the ``0-1\" test \\cite{GM04,GM09a,GM09b}. Interesting accounts of these methods can be found in \\cite{Con_spr}, as well as in a more recent review paper \\cite{SkoLE:2010}.  A novel, very efficient method based on the evolution of $k\\geq 2$ initially linearly independent deviation vectors is provided by the so-called ``Generalized ALignment Indices\" (GALI or GALI$_k$ spectrum) introduced in \\cite{SBA:2007} as a generalization of the ``Smaller ALingment Index\" (SALI) \\cite{Sk_sali:2001,SABV:2003a,SABV:2004}. The major advantage of the GALI method is that it follows the evolution of two or more deviation vectors and is thus able to extract more information about the complexity of the motion, yielding i.e.~the dimensionality of the invariant torus on which a regular orbit lies and predicting faster the chaotic nature of trajectories \\cite{SBA:2008,BouManChris:2010,ManRuf:2011}. To date, the GALI and the SALI indices have been successfully applied to a wide variety of autonomous (i.e.~explicitly time-independent) conservative flows and maps (see e.g.~ \\cite{PBS:2004,SESS,BouSko:2006,AntBou:2006,ABS:2006,CDLMV,VHC,KEV,MSAB:2008,Stra2009,BP09,Macek2010,SG10,MA11,BCSV12,MSA:2012,BCSPV12,GES12}). A concise review of the theory and applications of both the SALI and GALI methods can be found in \\cite[Chapter 5]{BS12}.  The motivation of the current work is twofold: First, we wish to investigate the different dynamical properties of a non-autonomous galactic potential, whose time-dependence could mimic certain realistic general trends arising in barred N-body galaxy simulations. Our second main goal is to explore the advantages of the GALI method, over the more traditional LEs, in detecting dynamical transitions in Hamiltonian systems, whose equations of motion are explicitly time-dependent.  There are, of course, several studies of time-dependent (TD) galactic and cosmological models in the literature, which use different tools to identify the chaotic vs.~regular nature of orbits. Defining the orbital complexity $n(k)$, of an orbital segment as the number of frequencies in its discrete Fourier spectrum that contain a $k$-fraction of its total power \\cite{KandEckBra:1997}, one may compare $n(k)$ with the short-time evolution of the LEs for TD models \\cite{Sio_etal:1998}. In \\cite{KandDru:98} the case of a cosmological model is discussed, where orbits may experience regular and/or chaotic motion during their time evolution, while in \\cite{SioKand:2000} the effects of a black hole, friction, noise and periodic driving are studied on a triaxial elliptic galaxy model, in which a type of transient chaos was found caused by a damped, oscillatory component \\cite{KandVassSid:2003,TerzKand:2004}. Finally, in \\cite{Sid:2009} the so-called ``pattern method\" was used to study a H\\'{e}non-Heiles potential to which an exponential function of time is added, while the dynamics of some simple TD galactic models was investigated in \\cite{CarPap:2003,Zot:2012}.   We recall here that in conservative systems the asymptotic nature of an orbit may be either periodic, quasiperiodic or chaotic. In the latter case, however, it may take a very long time before one can safely claim that a ``final\" state is reached, depending on the local dynamical properties, which may be characterized by ``strong or weak chaos\". In \\cite{Kats:2011a,Kats:2011b,MSA:2012} the dynamics in the vicinity of periodic orbits in conservative systems was studied by means of the maximum Lyapunov exponent (MLE) and the GALI. Here we explore the advantages of the GALI method and compare its predictions with what one finds using more traditional methods like the computation of the MLE also for TD systems.  In particular, we focus our attention on the dynamics of a barred galaxy model containing a disc and a bulge component, which is a widely accepted model for real barred galaxies. In the spirit of a mean field approach, we consider the motion of stars (represented by point particles) in this potential. The richness of the dynamics of the time-independent (TI) version of this model has been extensively studied in terms of: (a) the detection of periodic orbits and the analysis of their stability (see e.g.~\\cite{ABMP,Pfe:1984a,SPA02a,SPA02b,PSA02,PSA03a,PSA03b,SPPV05}), (b) the estimation of the relative fraction of chaotic vs.~regular orbits \\cite{MSAB:2008,MA:2009,MA11}, and (c) the statistical distributions of orbital coordinates described by $q$-Gaussian distribution functions \\cite{BouManAnt:2012}.  Here, we extend the analysis by considering a TD version of this model. More specifically, we allow some mass parameters of the potential to vary linearly as functions of time. As expected, whether we study a 2 dof or 3 dof version of the model, these variations can change the stability properties of periodic orbits, ``dissolve'' islands of regular motion and alter the structure of phase space in very complicated ways. Furthermore, in the TD case, the vast majority of dynamical transitions of phase space orbits cannot be claimed to be due to stickiness phenomena or ordinary diffusion to different regimes, as expected for TI Hamiltonian systems.  Recently, it was found in the TI case that the relative fraction of chaotic orbits grows as the bar's strength increases \\cite{MA11}. A question therefore arises, whether a similar correlation holds in the TD model, in the presence of ``realistic\" trends, which permit the mean field potential to vary in a way that is compatible with self-consistent N-body simulations regarding several components of the system.  Clearly, the analysis of the full N-body problem describes much better the galactic evolution and captures in great detail the different stellar structures present in the dynamics. However, there are serious difficulties and limitations when one tries to apply dynamical chaos detectors to such ``realistic\" many-particle systems due to the lack of sufficient orbital information during the time evolution. For this reason, many researchers prefer to use mean field potentials that are ``frozen\" in time and study the properties at specific snapshots of the simulations \\cite{Kalapo:2008,HaKa:2009}.  Keeping in mind that a barred galaxy experiences several dynamical transitions in different epochs that cannot be easily incorporated in our TD mean field potential, we shall proceed by making some helpful assumptions in an attempt to understand the behavior of such widely used chaos detectors as the GALI and the MLE. Thus, we will treat here two very general dynamical trends known to occur in barred galaxies: In the first scenario the mass of the bar component grows linearly in time (at the expense of the disc mass). This increase may be caused by an exchange of angular momentum with the disc (outer parts gain momentum from the inner parts), as has already been observed in N-body simulations (see e.g.~\\cite{AthMiss:2002,Ath:2003}). The fundamental trend in this case is that bars generally grow stronger in time, become more elongated and massive and eventually slow down.  We will also consider the inverse scenario, where the bar gets weaker making the disc more massive as time evolves (see e.g.~\\cite{Combes:2008,Combes:2011}).  The paper is organized as follows: In section \\ref{Model_gal} we present the TD barred galaxy model used in our study, while section \\ref{methods} is devoted to the description of the numerical methods employed for the computation of the MLE and the GALIs. Section \\ref{results} contains the main numerical results of the paper. A detailed investigation of the dynamics of particular orbits in a 2 dof version of our model is performed in section \\ref{2DOF}, while orbits of the full, 3 dof model, are considered in section \\ref{sec:3dof}. A global investigation of the dynamics of our TD galactic model is given in section \\ref{sec:global}. Finally, in section \\ref{sec:conclusions} the main conclusions of our work are presented. ",
        "conclusions": "\\label{sec:conclusions}  Autonomous Hamiltonian systems are conveniently studied for  fixed values of the total energy, where the location and extent of their regular and chaotic regions are time independent and can be accurately identified by a variety of methods especially in the low degree of freedom case. Even in such TI systems the dynamics can exhibit remarkable complexity, as there exist regimes of ``strong'' and ``weak'' chaos, as well as varying degrees of regularity, as the motion can occur on invariant tori of different dimensions and exhibit surprising localization properties in configuration and frequency space \\cite{BS12}.  Naturally, therefore, Hamiltonian systems which are explicitly TD are expected to be a lot more complicated, since, in the absence of TI integrals, all the above attributes evolve in time. For example, in the TI case, orbits do not change their nature: If they are initially regular they will always remain so, while if chaotic they can get trapped for long times on the boundary of regular regimes (exhibiting a ``weak'' form of chaos), but will never entirely relinquish their chaoticity. This is not so in the TD case, where individual trajectories may indeed display sudden transitions from regular to chaotic behavior and vice versa during their time evolution.  In the present paper, we have sought to shed some new light on these fascinating phenomena by studying the dynamics of a mean field model of a barred galaxy, whose mass parameters are allowed to vary linearly in time. Our primary goal was to show that transitions from order to chaos and vice versa do occur in this model and can be monitored much more accurately by local methods such as the GALI spectrum, rather than the more traditional approach of LEs. In addition, we wanted to investigate some astronomical properties of this TD system as it does incorporate some of the features appearing in N-body simulations and TI analytic potentials.  In this regard, we have chosen for simplicity to vary only two parameters in time, $M_B$ and $M_D$, keeping the size of the bar and the pattern speed fixed. Since the total mass is constant, whatever mass the bar loses is gained by the disc component and vice-versa. Moreover, to investigate more thoroughly the observed dynamical transitions, we have extended the maximal integration time of the orbits to $T=20000\\;Myr$ (20 billion years), which corresponds to a total interval of nearly 2 Hubble times. However, our results demonstrate that rich behavior and fundamental changes can also be observed within a single Hubble time of $10000\\;Myr$.  Most importantly, the GALI method turns to be again very efficient and accurate in the detection of chaotic motion in a TD system, as in the case of TI models. Furthermore, our work reveals that the method is especially suited for detecting intervals where an orbit changes its state fundamentally. By following the times that the GALIs require to fall to zero, one can describe in detail the orbit's successive passages from order to chaos and vice versa.  Finally, we focus on a more global astronomical study, where one would like to estimate qualitatively and quantitatively the relative fraction of regular and chaotic motion in such galaxy models. To this end, we choose different sets of initial conditions, launch them in phase space and classify regular and chaotic orbits, depending on whether the GALI fluctuates around a non-zero value or falls exponentially to zero. We have thus been able to verify that the conclusion of an earlier publication on the TI model \\cite{MA11}, that the percentage of chaos grows as the mass of the bar increases, holds true in the TD case as well. It would be highly interesting to investigate these questions in more realistic models, where besides the mass parameters the rotation frequency of the galaxy is also allowed to vary accordingly.  In closing, it is important to point out the advantages of the GALI method over the computation of the finite time MLE, as the GALIs do succeed in clearly capturing the transitions between different regular states and identifying the intermediate chaotic phases. By contrast, the manifestation of these different dynamical behaviors is much less pronounced in the time evolution of the MLE described by $\\sigma_1$ in this paper. Evidently, the practice of averaging Lyapunov exponents over an orbit's history smoothens out their fluctuations over short-lived events and gives them meaning only in the sense of the long time limit."
    },
    "1208/1208.0849_arXiv.txt": {
        "abstract": "We investigate magnetic reconnection and particle acceleration in relativistic pair plasmas with three-dimensional particle-in-cell (PIC) simulations of a kinetic-scale current sheet in a periodic geometry.  We include a guide field that introduces an inclination between the reconnecting field lines and explore outside-of-the-current sheet magnetizations that are significantly below those considered by other authors carrying out similar calculations. Thus our simulations probe the transitional regime in which the magnetic and plasma pressures are of the same order of magnitude.  The tearing instability is the dominant mode in the current sheet for all guide field strengths, while the linear kink mode is less important even without guide field.  Oblique modes seem to be suppressed entirely.  In its nonlinear evolution, the reconnection layer develops a network of interconnected and interacting magnetic flux ropes.   As smaller flux ropes merge into larger ones, the reconnection layer evolves toward a three-dimensional, disordered state in which the resulting flux rope segments contain magnetic substructure on plasma skin depth scales. Embedded in the flux ropes, we detect spatially and temporally intermittent sites of dissipation reflected in peaks in the parallel electric field. Magnetic dissipation and particle acceleration persist until the end of the simulations, with simulations with higher magnetization and lower guide field strength exhibiting greater and faster energy conversion and particle energization.  At the end of our largest simulation, the particle energy spectrum attains a tail extending to high Lorentz factors that is best modeled with a combination of two additional thermal components.  We confirm that the primary energization mechanism is acceleration by the electric field in the X-line region.  The highest energy positrons (electrons) are moderately beamed with median angles $\\sim 30^\\circ-40^\\circ$ relative to (the opposite of) the direction of the initial current density, but we speculate that reconnection in more highly magnetized plasmas would give rise to stronger beaming. Lastly, we discuss the implications of our results for macroscopic reconnection sites, and which of our results may be expected to hold in systems with higher magnetizations. ",
        "introduction": "Magnetic reconnection \\citep[e.g.,][and references therein]{yamada_magnetic_2010} is of interest in diverse areas of astrophysics, yet its mechanics remains incompletely understood.  A tearing instability is thought to be necessary to initiate reconnection in reversing magnetic field configurations, but many astrophysical plasmas have collisional resistivities that are insufficient to facilitate its growth. Interpretations of space plasma measurements \\citep[e.g.,][]{chen_observation_2008,oieroset_direct_2011} and astronomical observations, however, suggest that efficient collisionless reconnection is ubiquitous. Therefore, collisionless effects, which operate on plasma kinetic scales, seem to be required to provide the dissipation necessary for effecting a change of magnetic topology.  Magnetohydrodynamic (MHD) models for the dynamics of systems undergoing magnetic reconnection have been available for a long time \\citep[e.g.,][and references therein]{priest_magnetic_2000}.  Unfortunately, these models do not describe the underlying nature of (possibly multiscale) plasma organization in the reconnection layer where magnetic energy is being dissipated and the assumptions of ideal MHD do not apply.  Understanding the detailed plasma organization on all length scales, from the likely relatively small, plasma kinetic scales, to the potentially much larger scales on which astrophysical dynamical systems ``prepare'' reconnection sites, and where ideal MHD may be valid, is paramount for completing the theories of a wide variety of astrophysical phenomena and for interpreting space plasma measurements and astronomical observations.  In an effort to develop a picture of magnetic reconnection from first principles, recent particle-in-cell (PIC) simulations have examined the dynamics of reconnection layers that start with a current sheet as thin as the plasma skin depth.  PIC simulations of reconnection in pair plasmas have been carried out in two spatial dimensions \\citep{zenitani_generation_2001,zenitani_particle_2007,jaroschek_fast_2004,bessho_collisionless_2005,bessho_fast_2007,bessho_fast_2010,bessho_fast_2012,daughton_collisionless_2007,hesse_dissipation_2007,drake_magnetic_2010,hoshino_stochastic_2012} and three dimensions \\citep{zenitani_three-dimensional_2005,zenitani_role_2008,yin_three-dimensional_2008,liu_particle_2011,sironi_acceleration_2011}. Among these, several \\citep{zenitani_three-dimensional_2005,zenitani_role_2008,bessho_fast_2007,hesse_dissipation_2007,hoshino_stochastic_2012} have investigated the role of departure from the idealized, exactly antiparallel reconnection by introducing a perpendicular ``guide'' field.  These various simulations have revealed novel forms of small-scale plasma self-organization that are interesting in their own right, but that must ultimately be related to and embedded within the appropriate larger astrophysical contexts \\citep[e.g.,][]{uzdensky_fast_2010}. While spacecraft measurements, which can be done in situ, can provide direct clues how to establish this embedding in space plasmas, in extrasolar contexts only an indirect relation can be established between the reconnection process and the observed emission \\citep[e.g.,][]{sironi_acceleration_2011,cerutti_beaming_2012}.  Common features seen in many PIC simulations of magnetic reconnection include the formation of chains of magnetic flux ropes (in three dimensions with a guide field; otherwise, the common terms ``islands'' or ``plasmoids'' may still be more appropriate), the merging of smaller flux ropes into larger ones, and an energization of the plasma in the reconnection layer.  In three dimensional simulations, kink-like and oblique modes, as well as secondary instabilities, can impart three dimensional structure to the reconnection layer.  Typically, the simulations are initialized in the so-called Harris sheet equilibrium describing a current sheet with a thickness similar to the plasma skin depth. The tearing instability first sets in on scales of the initial current sheet thickness.  Its nonlinear development produces a chain of skin-depth-scale flux ropes alternating with magnetic X-lines, the three dimensional generalization of two dimensional X-points.  In X-lines, violation of flux freezing and magnetic line reconnection can be facilitated by a pressure tensor anisotropy \\citep[][see also, e.g., \\citealt{hesse_dissipation_2007}, and references therein]{vasyliunas_theoretical_1975}. Smaller flux ropes tend to merge with each other to form larger ones; this gives rise to magnetic organization on increasingly larger spatial scales.    Three-dimensional PIC simulations of guide field reconnection in electron-ion plasmas exhibit these same features, e.g., \\citet{daughton_role_2011} found that oblique modes dominated over tearing modes when guide field was strong. Additional effects specific to plasmas with electron-ion mass disparity have also been identified, but are not relevant for the present work.  Energization of particles in reconnection layers has been investigated in a number of PIC simulations \\citep{zenitani_generation_2001,zenitani_particle_2007,jaroschek_fast_2004,drake_electron_2006,drake_magnetic_2010,bessho_fast_2007,bessho_fast_2010,bessho_fast_2012,egedal_formation_2009,huang_mechanisms_2010,oka_electron_2010,liu_particle_2011,sironi_acceleration_2011,egedal_large-scale_2012,hoshino_stochastic_2012,cerutti_beaming_2012}. Less attention has been given to particle energization in the general case of reconnection with a guide field in three dimensions \\citep{zenitani_role_2008}.  In two dimensional simulations, particle acceleration producing a nonthermal energy spectrum, an apparent power-law, is often reported. \\citet{cerutti_beaming_2012}, however, instead detect a new ultrarelativistic thermal component energized by the reconnection.   That the reconnection layers should energize particles is in agreement with analytical considerations \\citep[e.g.,][]{speiser_particle_1965,larrabee_lepton_2003,giannios_uhecrs_2010,uzdensky_reconnection-powered_2011,cerutti_extreme_2012}, which find that particles in the vicinity of the X-line in the reconnection layer are accelerated by the nearly-uniform electric field as they repeatedly cross, and are trapped within the converging plasma flows.  Other mechanisms focusing on energetic particles that have moved from the X-line region into the flanking islands have also been suggested \\citep[e.g.,][]{drake_electron_2006,drake_magnetic_2010}.   In three dimensional simulations, evidence for a nonthermal spectrum is less solid. It remains poorly understood which processes limit the energy to which particles can be accelerated in fully dynamical, three-dimensional reconnection layers, and how do the particle energy spectrum, the degree of accelerated particle beaming, and the temporal evolution of the accelerated population depend on the parameters of the reconnection layer.  In this work we employ three dimensional PIC simulations to investigate the evolution of current sheets in relativistic pair plasmas undergoing magnetic reconnection.  Our simulations add to the small but growing family of three dimensional PIC simulation of relativistic reconnection with a guide field.   With the intention to complement existing work, we initialize our simulations slightly differently than it is normally done, not assuming the usual Harris sheet equilibrium.  Also, we explore a parameter regime, involving magnetic to kinetic pressure ratios of the order of unity, that has thus far not received sufficient attention.  We observe an evolution of magnetic field geometry that constrains the viability of models for high-Lundquist-number reconnection layers in which the diffusion region contains a hierarchy of interacting plasmoids \\citep[e.g.,][]{shibata_plasmoid-induced_reconnection_2001,uzdensky_fast_2010}. The simulations also allow us to explore the character of particle energization in dynamical, fully three dimensional reconnection layers.  The paper is organized as follows.  Section \\ref{sec:simulations} describes our methodology and simulation setup, while  Section \\ref{sec:results} presents the results.  Section \\ref{sec:discussion} discusses our findings concerning development of kinetic instabilities in the current sheet, as well as our findings on particle energization, in view of the existing work on these topics. Finally, Section \\ref{sec:conclusions} reviews our main conclusions. ",
        "conclusions": "\\label{sec:conclusions}  In this paper, we carried out three dimensional PIC simulations of magnetic reconnection in a relativistic pair plasma with varying guide field strength.  Plasma magnetizations, expressed in terms of the magnetic to kinetic pressure ratio, were of the order of unity.  The initial conditions differed from the usual Harris sheet configuration by not having a large density contrast between the center of the current sheet and the background plasma.  We investigated the growth of unstable kinetic modes in the current sheet, as well as the nonlinear development of a three dimensional flux rope network.  We also investigated the character and efficiency of particle energization.  Our main results can be summarized as follows:  The current sheets in all simulations develop significant magnetic reconnection accompanied with conversion of magnetic to particle kinetic energy.  With the aid of Fourier decomposition, we ascertained that the linear tearing mode is dominant in the early evolution of the current sheet.  We find that no significant growth occurs in the linear kink and oblique modes. The nonlinear development of the tearing mode produces a chain of flux ropes separated by primary X-lines.  The flux ropes merge in hierarchical fashion whereby the merging time scale is proportional to the flux rope separation. During this phase magnetic reconnection takes place at the X-lines. We find that the dimensionless reconnection rates $\\sim (0.05-0.08)$  and the maximum outflow speeds $\\approx 0.4\\, c \\sim v_{\\rm A}/2$ in our simulations are similar to those detected in other three dimensional simulations of reconnection in pair plasmas. We also find that spatial variation of an off-diagonal component of the pressure tensor is responsible for the breaking of flux freezing  at the X-lines, consistent with existing results.  While the hierarchical flux rope merging process initially appears similar to that found in two dimensional simulations, in fact it is three-dimensional from the outset.  This is because a lack of initial phase coherence in the linear tearing mode on scales larger than those allowed by causality breaks translational invariance in the direction of the initial current flow.  The flux ropes form a topologically interconnected, dynamically evolving network. Dynamical interaction between neighboring flux ropes is provided by magnetic tension forces. With time, the flux ropes break up into segments with more isotropic orientations.  The strongly three-dimensional character of the reconnection layer seems to suggest that global reconnection models invoking quasi-two-dimensional plasmoid hierarchies \\citep[e.g.,][]{shibata_plasmoid-induced_reconnection_2001,fermo_statistical_2010,uzdensky_fast_2010} require revision to account for the inter-plasmoid magnetic linkage and isotropization of plasmoid orientations.  The larger flux ropes produced during flux rope merging contain substructure down to plasma skin depth scales which is reflected in embedded, twisted and braided current filaments and sheets. Overall, this substructure is force-free and evolves relatively slowly. However, isolated sites within the evolved flux rope network contain spatially and temporally intermittent sites characterized by strong nonideal conditions ${\\mathbf E}\\cdot{\\mathbf B}\\neq 0$ where a change of magnetic connectivity continues to take place even after flux rope merging has saturated on length scales equal to the size of the computational box.  This intermittency may produce the observed variability of nonthermal emission in systems in which the emitting particles are energized by magnetic reconnection.  During the early, ordered flux rope merging phase, particles are accelerated to high Lorentz factors by the electric field in primary X-lines; the trajectories of these particles are well described by Speiser orbits. Particles continue to be energized in the later, disordered phase we identify in our largest simulation, but we leave the analysis of energization in the disordered regime to a subsequent investigation.  Simulations with higher magnetization and lower guide field strength exhibit greater and faster energy conversion and particle energization.  The efficiency of particle energization measured in terms of the energy in the accelerated particles per unit magnetic energy dissipated in the simulation is an increasing function of the guide field strength, which can be interpreted as resulting from a decreasing plasma compressibility with increasing guide field. The final particle energy spectrum in the largest simulation is best fit by the inclusion of new thermal components at temperatures $2.1\\,m_e c^2$, and $3.5\\,m_e c^2$, in addition to the initial thermal component with temperature $m_e \\,c^2$.  We, however, acknowledge that a larger size or longer duration simulation is likely to produce a still more pronounced energized component, possibly even a population described with a power law spectrum.  Energetic positrons (electrons) with Lorentz factors $\\gamma > 30$ are moderately beamed in (opposite to) the direction of the initial current flow with median inclinations of $\\sim30^\\circ-40^\\circ$. The degree of beaming is determined by a particle's energy gain during acceleration.  We speculate that more highly magnetized plasmas and reconnection sites with larger size X-line regions should give rise to stronger beaming.  In this work, we have investigated a narrow range of magnetizations with $\\sigma\\sim {\\mathcal O} (1)$, but astrophysical reconnection sites can also have high magnetizations $\\sigma\\gg 1$.  We can speculate about the applicability of our results in the latter limit. The linear tearing mode responsible for the initiation of reconnection is insensitive to the degree of magnetization far from the current sheet.  The phase decoherence that produces the initial breakdown of translational invariance is determined by the tearing mode growth time and should thus also persist at high magnetizations. Therefore, we expect the qualitative structure of the reconnection region at higher values of $\\sigma$ to be similar to that found in our simulations.  The primary effect of high magnetization is that the Alfv\\'en velocity approaches the speed of light, which could give rise to ultrarelativistic outflows from the X-line region. In such outflows the inertial term of the generalized Ohm's law becomes important in the breaking of flux freezing \\citep{hesse_dissipation_2007}. This in turn may increase the dimensionless reconnection rate $r_{\\rm rec}$ relative to the value found in our simulations.  It remains to be seen whether the associated reconnection process is more or less intermittent.  An increased magnetization is likely to increase the efficiency and the degree of beaming in particle energization.  We consider these results and the immediate questions they raise an incremental step in the development of a multiscale view of collisionless plasma self-organization during magnetic reconnection.  Further work is clearly required to place our key finding, the evolution of the simulated, periodic reconnection layer into a disordered network of interacting magnetic flux ropes, in the macroscopic context of a realistic reconnection site characterized by outflow boundary conditions and altogether different field line asymptotics at large distances from the X-line.   It will be particularly interesting to see if the reconnected-flux-carrying outflow from the macroscopic X-line will possess the disordered, interlinked magnetic field topology we observe and investigate what will be the character of magnetic fluctuations in the outflow."
    },
    "1208/1208.5836_arXiv.txt": {
        "abstract": "We present a comprehensive data description for Ks-band measurements of Sgr~A*. We characterize the statistical properties of the variability of Sgr~A* in the near-infrared, which we find to be consistent with a single-state process forming a power-law distribution of the flux density. We discover a linear rms-flux relation for the flux-density range up to 12 mJy on a timescale of 24 minutes. This and the power-law flux density distribution implies a phenomenological, formally non-linear statistical variability model with which we can simulate the observed variability and extrapolate its behavior to higher flux levels and longer timescales. We present reasons why data with our cadence cannot be used to decide on the question whether the power spectral density of the underlying random process shows more structure at timescales between 25 min and 100 min compared to what is expected from a red noise random process. ",
        "introduction": "\\label{intro}  Sagittarius (Sgr~A*) at the center of the our galaxy is a highly variable near-infrared (NIR) and X-ray source which is associated with a $ 4 \\times 10^{6} M_{\\sun}$ supermassive central black hole (\\citealt{1996Natur.383..415E,1997MNRAS.284..576E}; \\citealt{2002MNRAS.331..917E}; \\citealt{2002Natur.419..694S}; \\citealt{2003ApJ...597L.121E}; \\citealt{1998ApJ...509..678G,2000Natur.407..349G,2005ApJ...620..744G,2008ApJ...689.1044G}; \\citealt{2009ApJ...692.1075G}). While first detected as a bright, ultra compact, and comparatively steady radio source, the strong variability at shorter wave-lengths, the variable polarization of the NIR emission, and the correlation between fluctuations in the the sub-mm, NIR and X-ray regimes provide evidence that this variable emission originates in the direct surrounding of the black hole. Therefore, properties of the black hole and of the emission and accretion mechanisms in its close surrounding can be studied at these wavelengths (\\citealt{2001Natur.413...45B}; \\citealt{2003A&A...407L..17P,2008A&A...488..549P}; \\citealt{2003Natur.425..934G}; \\citealt{2004ApJ...601L.159G,2005ApJ...635.1087G} ;\\citealt{2004A&A...427....1E, 2006A&A...450..535E, 2006A&A...455....1E, 2006Msngr.125....2E, 2008A&A...479..625E, 2008JPhCS.131a2002E, 2008A&A...492..337E}; \\citealt{2006ApJ...644..198Y, 2006ApJ...650..189Y, 2007ApJ...668L..47Y, 2008ApJ...682..361Y}; \\citealt{2007ApJ...667..900H} \\citealt{2009ApJ...698..676D}; \\citealt{2010A&A...512A...2S}).  To explain the observed variability and its correlation between the NIR and X-ray regimes several authors suggest Synchrotron Self-Compton (SSC) or inverse Compton emission as the responsible radiation mechanisms (\\citealt{2004A&A...427....1E, 2006A&A...450..535E, 2006A&A...455....1E}; \\citealt{2004ApJ...606..894Y}; \\citealt{2006ApJ...648.1020L}; \\citealt{2012A&A...537A..52E}). Relativistic models that assume the variability to be linked to emission from single or multiple hot spots in the accretion disk near the last stable orbit of the black hole have been applied successfully to individual datasets (\\citealt{2006A&A...460...15M, 2006A&A...458L..25M, 2007A&A...473..707M}; \\citealt{2008JPhCS.131a2008Z}). These models interpret the shorter timescales of the variability (between 10 and 30 minutes) to be dominated by the timescales of the orbital motion\\footnote{$ \\sim 20 $~minutes at the innermost stable orbit for Sgr~A*.}. Orbital motion close to the black hole (and an associated quasi periodic signal in the light curves) is of special interest: it could used as a timing experiment in a strong gravitational field that might allow for determining black hole parameters like spin or inclination.  Potential quasi-periodic oscillations (QPOs) have been found in some light curves (see, e.g., \\citealt{2003Natur.425..934G}). However, these QPOs are not statistically significant within the overall variability. Based on NIR light curves of 7 nights observed with Keck telescope \\cite{2009ApJ...691.1021D} analyzed the flux density distribution of Sgr~A* and the significance of quasi periodic oscillations. They showed that QPOs in total intensity light curves cannot be established with sufficient significance against random fluctuations, finding a pure red noise power spectral density sufficient to account for the time correlation of the fluctuations.  On the other hand, based on relativistic models, \\cite{2010A&A...510A...3Z} predicted a correlation between the modulations of the observed flux density light curves and changes in polarimetric data (also see \\citealt{2006A&A...450..535E}, \\citealt{2006A&A...460...15M, 2006A&A...458L..25M, 2007A&A...473..707M}; \\citealt{2008A&A...479..625E}, \\citealt{1973ApJ...183..237C}; \\citealt{1977Natur.266..429S}; \\citealt{1991A&A...245..454A}; \\citealt{1992A&A...257..531K}; \\citealt{1995ApJ...448L..21H}; \\citealt{2004ApJS..153..205D,2008MNRAS.384..361D}). A comparison of predicted and observed light curve features (obtained from 6 nights of polarimetric observations with VLT and Subaru telescope) through a pattern recognition algorithm resulted in the detection of a signature possibly associated with orbiting matter under the influence of strong gravity.  Since the first discovery of variable emission of Sgr~A* in the NIR in 2003 (\\citealt{2003Natur.425..934G}) a number of publications concentrated on the statistical properties of the flaring activity rather than on interpreting individual observations. These papers investigated the timing properties of the light curves as well as the radiation mechanisms involved (\\citealt{2009ApJ...706..348Y}, \\citealt{2010A&A...510A...3Z}, \\citealt{2011A&A...532A..26B}, \\citealt{2011A&A...532A..83S}, \\citealt{2012A&A...537A..52E}). On the base of 14 light curves observed between 2004 and 2008, including alternating observations with VLT and Keck, \\cite{2009ApJ...694L..87M} discovered first a dominant timescale at about 150 min, supporting linear scaling relations of break timescales in the power spectral density with the black hole mass. They determined the power-law slope of the high-frequency part of the power spectral density (PSD) to be $ 2.1 \\pm 0.5 $.  A comprehensive statistical approach in the analysis of the Ks-band total intensity variability observed with NACO at the VLT has been conducted by \\cite{2011ApJ...728...37D}. The authors analyzed VLT K-band data between 2004 and 2009. They presented a detailed investigation of the flux density statistics and described the time-variable stellar confusion at the position of Sgr~A* that makes an investigation of the faint emission difficult. They also emphasized the importance of these faint states for the overall statistical evaluation of the variability of Sgr~A*.  Based on the flux density histogram the authors claim the evidence for two states of variability, a log-normal distributed quiescent state for flux densities $ <5 $ mJy and a power-law distributed flaring state for flux densities $ >5 $ mJy. They argue that it is very unlikely that the same variability process is responsible for both high and low flux density emission from Sgr~A* (the statistical model the authors promote is summarized in Appendix~\\ref{domo}). With reference to this model \\cite{2010RvMP...82.3121G} state that ``it is a key issue whether the brightest variable emission from Sgr A* are statistical fluctuations from the probability distribution at low flux or flare events with distinct properties\" and that ``the transition from the log-normal distribution at low-flux levels to the tail of high fluxes may also explain the apparent mismatch between the detection vs. non- detection of quasi-periodic substructures in different near-infrared light curve studies\".  Our statistical analysis presented in this paper serves the following goals:  \\begin{itemize} \\item to provide a more comprehensive, uniformly reduced data set of Ks-band observations from 2003 to early 2010;  \\item to conduct a rigorous analysis of the observed flux density distribution;  \\item to explain why a proper statistical analysis of the Ks-band light curves cannot reproduce the results found by \\cite{2011ApJ...728...37D};  \\item to conduct a rigorous time series analysis on the base of a representative dataset;  \\item  to propose a comprehensive statistical model that, using standard methods for generating Fourier transform based surrogate data, describes all aspects of the observed (total intensity) data and lets us simulate light curves with the observed time behavior and flux density distribution;  \\item and to investigate extreme variability events in the context of our statistics. \\end{itemize} ",
        "conclusions": "\\label{concl}  We summarize our results:  \\begin{itemize} \\item The Ks-band flux density distribution of Sgr~A* as obtained from the last seven years of observations is in convincing agreement with a pure power-law distribution, giving no indication for a break or two-state behavior.  \\item We could find an upper limit of the intrinsic mean flux density of about $ 1.7 \\pm 0.15 $ mJy, and with a power-law extrapolation to flux densities below the detection limit a mean of $ 0.9 \\pm 0.15 $ mJy (including extinction correction).  \\item We found an algorithm to statistically simulate light curves that show the same flux density distribution and time correlation as the observed sample. It is based on the algorithm by \\cite{1995A&A...300..707T} to generate linearly time-correlated surrogate samples, but includes a transformation to account for the non-linear appearance of the NIR flux densities of Sgr~A*. In a first approximation broken power-law PSDs are invariant under this transformation. This statistical model (best fitting PSD, flux density distribution and the algorithm) does not provide immediate information on the physical system, but serves as a statistical ``summary\" of the observed variability, defining constrains for every physical model. Furthermore, it allows a straight forward power-law extrapolation to higher flux density levels, flux density levels below the detection limit, and long timescales.  \\item This extrapolation demonstrates that high flux density excursions, as required to explain the supposed X-ray light echo in the molecular clouds surrounding Sgr~A*, are well within the expected statistical extreme values of the flux density distribution that we observed at much lower flux densities. This implies that even if we include the flare responsible for the X-ray echo as the most extreme event that is suggested by observations there is no evidence for a two-state variability behavior.  \\item Within our statistical model the concept of a ``quiescent state\" and the differentiation between continuous variability and ``off\" states (as investigated by \\citealt{2011ApJ...728...37D}) turn out to be problematic: In this description, Sgr~A* is always variable and the probability to find it at a flux density level of zero is actually zero. But any flux density interval starting with zero is represented with a higher probability than any other interval of the same length, allowing for arbitrarily faint flux density states. So from an observational point of view, this model predicts ``off\" states due to the instrument dependent limited resolution and sensitivity. In the case of Ks-band observations with NACO the detection limit is very close to the mean flux density of the intrinsic distribution.  This means that the distribution of about half of the intrinsic flux density states can not be accessed with our current resolutions and sensitivity. Next generation instruments like LINC-NIRVANA at the Large Binocular Telescope and GRAVITY at the VLTI will allow to probe the lower flux density range.  \\item The question whether timescales comparable to the orbital timescale at the innermost stable orbit play a role for the variability is principally undecidable on the base of this data. In order to access this problem a significant amount of continuous light curves with a length of more than 1500 min would be needed. This, however, does not exclude the presence of orbital signatures in polarimetric data as reported in \\cite{2010A&A...510A...3Z}. Here we are investigating all light curves in total under the assumption of stationarity. Thus, we cannot make statements on a possible time development of the dominant timescale or the role of shorter timescales as a transient phenomenon.  \\end{itemize}"
    },
    "1208/1208.4600_arXiv.txt": {
        "abstract": "{We discuss decays of ultra-relativistic neutrinos over cosmological distances by solving the decay equation in terms of its redshift dependence. We demonstrate that there are significant conceptual differences compared to more simplified  treatments of neutrino decay. For instance, the maximum distance the neutrinos have traveled is limited by the Hubble length, which means that the common belief that longer neutrino lifetimes can be probed by longer distances does not apply. As a consequence, the neutrino lifetime limit from supernova 1987A cannot be exceeded by high-energy astrophysical neutrinos. We discuss the implications for neutrino spectra and flavor ratios from gamma-ray bursts as one example of extragalactic sources, using up-to-date neutrino flux predictions. If the observation of SN 1987A implies that $\\nu_1$ is stable and the other mass eigenstates decay with rates much smaller than their current bounds, the muon track rate can be substantially suppressed compared to the cascade rate in the region IceCube is most sensitive to. In this scenario, no gamma-ray burst neutrinos may be found using muon tracks even with the full scale experiment, whereas reliable information on high-energy astrophysical sources can only be obtained from cascade measurements. As another consequence, the recently observed two cascade event candidates at PeV energies will not be accompanied by corresponding muon tracks.} ",
        "introduction": "While the masses and mixings of the neutrinos are already well probed, other  neutrino properties are less understood, such as the  electromagnetic properties and the neutrino lifetime.   The most stringent phenomenological bound on neutrino lifetime comes from the observation of neutrinos from supernova 1987A~\\cite{Hirata:1987hu,Bionta:1987qt}, which were measured in the electron flavor. Given the uncertainty on the supernova neutrino flux (order 50\\%) and the neutrino mixing parameters, it may apply to the mass eigenstate $\\nu_1$ or $\\nu_2$. For the sake of simplicity, we assume that the bound applies to $\\nu_1$: $\\tau_1/m_1 \\gtrsim 10^5~\\text{s/eV}$.\\footnote{Neutrino lifetime is usually described by  $\\kappa_i^{-1} \\equiv \\tau_{i,0}/m_i$, where $\\tau_{i,0}$ is the rest frame lifetime of the  mass eigenstate $\\nu_i$. The origin of the mass dependence is the fact that decays scale as $\\exp[-t / (\\tau_{i,0} \\gamma)] \\simeq \\exp[-(L m_i) / (E \\tau_{i,0})]$, \\ie, the rest frame lifetime $\\tau_{i,0}$ is  boosted by $\\gamma = E/m_i$ into the observer's frame, and $L$ (baseline source-detector) and $E$ (neutrino energy) are quantities related to source and experiment.} The (model-independent) bounds on the other mass eigenstates are less stringent: bounds on $\\nu_2$ lifetime are imposed by solar neutrino data, yielding $\\tau_2/m_2 \\gtrsim 10^{-4}~\\text{s/eV}$ for decays into invisible daughters~\\cite{Joshipura:2002fb,Bandyopadhyay:2002qg,Beacom:2002cb} and $\\tau_2/m_2 \\gtrsim 10^{-3}~\\text{s/eV}$ for decay modes with secondary $\\bar{\\nu}_e$ appearance~\\cite{Eguchi:2003gg, Aharmim:2004uf}. Furthermore, $\\nu_3$ is constrained from the analysis of atmospheric and long-baseline neutrino data, $\\tau_3/m_3 \\gtrsim 10^{-10}~\\text{s/eV}$~\\cite{GonzalezGarcia:2008ru}. More stringent bounds can be derived  when specific decay models are assumed (see, \\eg, \\Refs~\\cite{Pakvasa:2003db, Yao:2006px, Pakvasa:2008nx} for an overview). For example, solar neutrinos strongly limit the possibility of radiative decays~\\cite{Raffelt:1985rj}, while for Majoron decays~\\cite{Gelmini:1980re, Chikashige:1980qk}, explicit bounds can be obtained from neutrino-less double-beta decay and supernovae~\\cite{Tomas:2001dh}. Another possibility is the decay into un-particles~\\cite{Chen:2007zy, Zhou:2007zq, Li:2007kj, Majumdar:2007mp}. In this study, we do not consider specific decay models, but focus on the phenomenology of neutrino decay, given the bounds on the lifetimes above. Especially neutrino telescopes~\\cite{Aslanides:1999vq, Ahrens:2002dv, Tzamarias:2003wd, Piattelli:2005hz} are sensitive to neutrinos with an average energy and traveled distance many orders of magnitude larger than present neutrino experiments, and may be an interesting approach to probe neutrino decay.  Neutrino oscillations are typically assumed to be averaged out over cosmological distances (see \\Ref~\\cite{Farzan:2008eg} for a discussion), and the usual flavor mixing remains to be taken into account. As far as decays are concerned, one distinguishes between decays into products {\\em invisible} to the detector, such as sterile neutrinos, un-particle states, Majorons, or active neutrinos strongly degraded in energy, and decays into {\\em visible} states, \\ie, active neutrino flavors. The decays can be {\\em complete}, \\ie, all unstable mass eigenstates have decayed (see, \\eg, \\Ref~\\cite{Beacom:2002vi}), or {\\em incomplete}, \\ie, the decay spectral signature will be visible. A complete classification of complete (visible or invisible) decay scenarios has been performed in \\Ref~\\cite{Maltoni:2008jr}, while incomplete invisible decay has, for instance, been studied for active galactic nuclei (AGNs) in \\Ref~\\cite{Bhattacharya:2009tx,Bhattacharya:2010xj}. The description of incomplete visible decays is, in general, more complicated~\\cite{Lindner:2001fx, Lindner:2001th}, which is why we focus on incomplete invisible decay in this study (IC-40 refers to the 40-string configuration of the IceCube detector). The characteristic energy dependence of neutrino decay may lead to easily observable imprints in the neutrino fluxes in that case~\\cite{Bhattacharya:2009tx,Bhattacharya:2010xj}. On the other hand, it has been proposed to use the flavor composition at the detector, such as the ratio between muon tracks and cascades, to distinguish different decay scenarios~\\cite{Beacom:2002vi, Lipari:2007su, Majumdar:2007mp, Maltoni:2008jr,Bustamante:2010nq,Mehta:2011qb}. We will discuss both options in this study, where our main focus are the fluxes.  It is often believed that decays are always complete if the sources lie far enough away, which, in turn, allows for the test of very long neutrino lifetimes of the order of \\begin{equation} \\kappa^{-1} \\, \\left[ \\frac{\\vphantom{M} \\mathrm{s}}{\\mathrm{eV}} \\right] \\equiv \\frac{\\tau \\, [\\mathrm{s}]}{m \\, [\\mathrm{eV}]} \\simeq 10^2 \\, \\frac{L \\, [\\mathrm{Mpc}]}{E \\, [\\mathrm{TeV}]} \\, . \\label{equ:est} \\end{equation} As a consequence, especially objects with high redshifts may be well suited to test neutrino lifetime, and may potentially lead to bounds stronger than the one from SN 1987A. An example, which we are going to use in this work, are gamma-ray bursts (GRBs), with observed redshifts as high as $z \\simeq 6-8$; see \\Ref~\\cite{Kistler:2009mv} for the redshift and luminosity distribution. This leads to potentially strong constraints via \\equ{est}. However, note that the relationship between distance $L$ and redshift $z$ depends for $z \\gtrsim 0.1$ (or $L \\gtrsim 360 \\, \\mathrm{Mpc}$) on the cosmological distance measure to be used, which we are going to address in this study; see \\Refs~\\cite{Wagner:2011uy,Weiler:1994hw} for a discussion in the context of neutrino oscillations, and \\Refs~\\cite{Beacom:2003eu,Esmaili:2012he} in the context of pseudo-Dirac neutrinos. For the neutrino flux, especially redshifts $z \\simeq 1$ will dominate~\\cite{Baerwald:2011ee}, which is a consequence of the convolution of the star formation rate (including some redshift evolution function) and the contribution to the total flux scaling as $1/d_L^2$, where $d_L$ is the luminosity distance. This value is already significantly beyond the indicated $z \\simeq 0.1$, which means that cosmological effects have to be taken into account.  GRBs are not only an interesting class of candidate neutrino sources because they may originate from high redshifts, but also because the non-observation of neutrinos from GRBs~\\cite{Abbasi:2012zw} has started to constrain the gamma-ray--neutrino connection in conventional internal shock models~\\cite{Guetta:2003wi,Abbasi:2009ig}, based on the ideas in \\Ref~\\cite{Waxman:1997ti}. A re-calculation of the predicted neutrino fluence from the gamma-ray fluence has yielded an order of magnitude lower result~\\cite{Hummer:2011ms} for the same burst parameters (see also \\Ref~\\cite{Li:2011ah}), which is a possible explanation soon going to be tested by new data. We will use this new nominal prediction for the quasi-diffuse flux~\\cite{Hummer:2011ms}, calculated with the IC-40 bursts~\\cite{Abbasi:2011qc}, in this study. Note that model-specific~\\cite{Hummer:2011ms,Barranco:2012xj} and more generic~\\cite{He:2012tq} astrophysical uncertainties may imply even lower neutrino flux predictions. On the other hand, another plausible explanation for why nothing has been seen yet may be that neutrinos have actually (at least partially) decayed between source and detector. We will investigate the impact of this hypothesis on the predicted fluxes.  One of the very recent  puzzles of neutrino astronomy is the potential observation of two cascade events at IceCube at  PeV energies~\\cite{NeutrinoTalk}. In particular, assuming that these are of extragalactic origin, an interesting question is: Why have these been observed as cascades, and no muon tracks have been seen? This question may become more prominent with increasing statistics. In this work, we demonstrate that neutrino decay can provide an answer to that question, given the current constraints on neutrino lifetime; for alternative explanations, see \\Refs~\\cite{Barger:2012mz, Fargion:2012zc}. We will demonstrate that muon tracks may be strongly suppressed in the presence of neutrino decay compared to (especially electromagnetic) cascades, and that cascade measurements are much more powerful to find astrophysical neutrinos in this case. We will use GRBs as an example, but our conclusions can be applied to AGNs as well.  This study is organized as follows: We describe the decay framework in \\Sec~\\ref{sec:decay}, which consists of the general framework including cosmological distances, a simplified solution, and the proper solution of the redshift-dependent decay equation. We then discuss the implications for GRBs in \\Sec~\\ref{sec:grb}, as one possible example of extragalactic neutrino sources. In that section, we introduce several benchmark GRBs and the analysis techniques, discuss the impact of decay on the neutrino fluxes, and introduce a phenomenologically viable model for which we show the implications for fluxes and flavor composition. In \\Sec~\\ref{sec:icecube} we apply this model to a realistic stacking analysis of IceCube, and compare the predicted neutrino fluxes to current and future bounds. Finally, we conclude in \\Sec~\\ref{sec:summary}. Note that, while we use GRBs as test case, the results from this study apply to the neutrino propagation from other extragalactic sources as well, including cosmogenic neutrinos. ",
        "conclusions": "\\label{sec:summary}  We have discussed neutrino decays of ultra-relativistic neutrinos over cosmological distances, \\ie, for redshifts $z \\gtrsim 1$. These neutrinos may originate from extragalactic cosmic accelerators, such as AGNs or GRBs. As class of  decay models, we have considered decays into secondaries invisible to the detector, such as Majorons, unparticles, sterile neutrinos, or neutrinos strongly degraded in energy. The decays have been allowed to be incomplete, \\ie, the spectral signature of the decay may be seen.  We have demonstrated that simplified approaches to neutrino decay cannot be used over such distances, and that the differential decay equation has to be carefully solved as a function of redshift. We have also given a simple parameterization for the solution, which can be easily used. From a comparison of different approaches, we have demonstrated that conventional approaches to neutrino decay can be safely used for $z \\lesssim 0.1$ (or distances $L \\lesssim 360 \\, \\mathrm{Mpc}$), whereas the light-travel (or look-back) distance triggers the clock for neutrino decays over cosmological distances. As a consequence, the horizon is limited by the Hubble length $\\simeq 3.89 \\, \\mathrm{Gpc}$, which is the maximal distance scale testable by decay. This has several interesting consequences: First of all, the common notion ``the longer the distance, the longer lifetimes can be tested'' does not hold anymore. Second, the concept of complete decays (all neutrinos have decayed) cannot be assigned to long enough distances, but only to low enough energies, \\ie, it is an energy-dependent and not distance-dependent statement. And third, the neutrino lifetime bound from SN 1987A, which emitted neutrinos at much lower energies, cannot be exceeded by high-energy astrophysical neutrinos from AGNs and GRBs tested in neutrino telescopes. The same argument applies to the cosmogenic neutrinos from cosmic microwave (and infrared) background interactions, which are at even higher energies. Note that the SN 1987A neutrinos have only been measured in the  electron flavor, which means that this limit does not apply to all mass eigenstates. However, it is expected that the neutrinos from any future galactic supernova explosion will be measured in different flavors and neutral currents, which means that in long terms the best lifetime bounds are expected to come from supernova neutrinos.  While our results can be easily applied to AGNs as well, we have used GRBs as test case. Because of the SN 1987A constraint, we have been forced to choose $\\nu_1$ stable and $\\nu_2$ and $\\nu_3$ decaying with equal rates.\\footnote{We have demonstrated that saturating the SN 1987A bound for the $\\nu_1$ lifetime does not change the results, which means that it is equivalent to the studied model for the purpose considered.} We have shown that, safely within the current bounds for $\\nu_2$ and $\\nu_3$, the muon neutrino flux can be substantially reduced  by about one order of magnitude, whereas the electron neutrino flux is hardly affected. This has very interesting implications for neutrino telescopes: First, improved bounds on the neutrino flux from muon tracks can be interpreted in terms of the astrophysical parameters or the possibility that neutrinos are unstable. On the other hand, the electromagnetic cascade bounds are hardly affected by the neutrino lifetime, which means that reliable conclusions on the astrophysical objects can only come from cascades, and highlights the importance of dedicated cascade analyses for GRBs and other classes of objects. This would be different, of course, if muon tracks from astrophysical neutrinos were actually observed, since their discovery would constrain this possibility.  For GRBs, in particular, we have used a realistic sample of IceCube GRBs and the state-of-the-art technology  for the prediction of the quasi-diffuse neutrino flux to illustrate the consequences. We have demonstrated that, for the discussed decay scenarios and our nominal flux prediction using the same parameters as IceCube, even the full scale IceCube experiment may not find neutrinos from GRBs in muon tracks after ten years, whereas the perspectives for a detection in cascades are actually very good.  Finally, note that recently two cascade neutrino event candidates at PeV energies have been found.  Our scenario provides a plausible explanation why only cascades are seen, whatever the origin of the neutrinos is, whereas there are no accompanying muon tracks in spite of comparable effective areas. For GRBs, in particular, the PeV energies are exactly where we predict the maximum of the quasi-diffuse flux; see \\figu{icecube}, right panel. On the other hand, for the relevant search time window, we only expect 0.07 electromagnetic cascade events for the nominal (``no decay'') prediction from GRBs. The observed two events might in this case come from a strong statistical fluctuation, or a significant deviation of astrophysical parameters from their assumed mean values."
    },
    "1208/1208.1739_arXiv.txt": {
        "abstract": "We present a study of the evolution of several classes of {\\MgII} absorbers, and their corresponding {\\FeII} absorption, over a large fraction of cosmic history: 2.3 to 8.7 Gyrs from the Big Bang. Our sample consists of 87 strong ($W_r(\\MgII)>$~0.3 \\AA) {\\MgII} absorbers, with redshifts 0.2~$ < z <$~2.5, measured in 81 quasar spectra obtained from the Very Large Telescope (VLT) / Ultraviolet and Visual Echelle Spectrograph (UVES) archives of high-resolution spectra (R $\\sim$ 45,000). No evolutionary trend in $W_r(\\FeII)/W_r(\\MgII)$ is found for moderately strong {\\MgII} absorbers (0.3~$<W_r(\\MgII)<$~1.0 {\\AA}). However, at lower redshifts we find an absence of very strong {\\MgII} absorbers (those with $W_r(\\MgII)>$~1 {\\AA}) with small ratios of equivalent widths of {\\FeII} to {\\MgII}.  At high redshifts, very strong {\\MgII} absorbers with both small and large $W_r(\\FeII)/W_r(\\MgII)$ values are present. We compare our findings to a sample of 100 weak {\\MgII} absorbers ($W_r(\\MgII) <$~0.3 \\AA) found in the same quasar spectra by Narayanan et~al. (2007).  The main effect driving the evolution of very strong {\\MgII} systems is the difference between the kinematic profiles at low and high redshift. At high redshift, we observe that, among the very strong {\\MgII} absorbers, all of the systems with small ratios of $W_r(\\FeII)/W_r(\\MgII)$ have relatively large velocity spreads, resulting in less saturated profiles. At low redshift, such  kinematically spread systems are absent, and both {\\FeII} and {\\MgII} are saturated, leading to $W_r(\\FeII)/W_r(\\MgII)$ values that are all close to 1. The high redshift, small $W_r(\\FeII)/W_r(\\MgII)$ systems could correspond to sub-DLA systems, many of which have large velocity spreads and are possibly linked to superwinds in star forming galaxies.  In addition to the change in saturation due to kinematic evolution, the smaller $W_r(\\FeII)/W_r(\\MgII)$ values could be due to a lower abundance of {\\Fe} at high redshifts, which would indicate relatively early stages of star formation in those environments. ",
        "introduction": "\\label{sec:1}  Quasar absorption lines allow the study of galaxies and their halos with no bias toward specific environments because of the random distribution of lines of sight in the sky. In particular, high resolution studies of quasar absorption lines help us better characterize the kinematic properties and chemical content of the absorbing gas. The kinematic composition of individual galaxies can be studied statistically with quasar absorption lines (\\citealt{Charlton98}). Profile shapes of the absorption lines show clouds of gas at different velocities relative to galaxy centroids, establishing, for example, whether the absorption is consistent with absorption in the disk and/or the extended halo (\\citealt{Steidel02b}; \\citealt{Kacprzak11}). Also, understanding the nature and evolution of galaxies requires assessing their star formation history, which can be traced by their metal content (e.g., \\citealt{Matteucci08} and references therein). Although the number of massive stars is smaller than the number of less massive stars, the former have a larger impact on the galaxy's chemical enrichment through supernovae explosions. Different types of supernovae (SNe) contribute differently to this enrichment: core-collapse SNe (Type II) produce a larger contribution of $\\alpha$--elements (such as \\Mg) relative to \\Fe\\, within the first billion years of the formation of a stellar population. \\Fe\\, is mostly generated in a later phase by Type Ia SNe. Thus, the ratio of \\Fe\\, relative to \\Mg\\, may be used as a measure of the stage in the star formation history of a galaxy. Quasar absorption line studies are a probe of this kind of information.  Among the absorption features that can be used in extragalactic studies, the resonance doublet transitions of one $\\alpha$--element, \\MgII\\,\\lb\\lb2796,2803, are well suited due to their strength and their accessibility at optical wavelengths from redshifts of 0.2 up to 2.7, covering a large range in cosmic ages (from approx. 2.5 to 11.3 Gyrs). These intervening {\\MgII} absorption lines have been historically (and originally arbitrarily) classified based on the strength of their rest-frame equivalent width in the 2796~{\\AA} transition ($W_r({\\MgII})$\\footnote{Throughout the paper we use $W_r(\\MgII)$ to describe the rest-frame equivalent width of the {\\MgII} \\lb2796 transition}). Posteriorly, physical properties and origins for these systems have defined statistically distinct populations, although the equivalent width boundaries between these different populations are not sharp.  Strong {\\MgII} absorption systems are those with $W_r({\\MgII}) \\geq$~0.3 \\AA, while weak \\MgII\\, absorption is defined to have $W_r({\\MgII}) <$~0.3 \\AA. Several sub-classifications have been proposed among the strong absorption systems: very strong absorbers are those with $W_r({\\MgII}) \\geq$~1 \\AA, and the term ``ultrastrong'' is reserved for those \\MgII\\, absorption systems with $W_r({\\MgII}) \\gsim$~3 \\AA\\, (\\citealt{Nestor05}; \\citealt{Nestor07}).  Strong \\MgII\\, absorption has generally been found to be connected to galaxies. Luminous galaxies ($\\sim0.1-5{L_\\star}$ galaxy) near quasar lines of sight show strong \\MgII\\, absorption within 60~kpc in $\\approx$~75\\% of all cases (\\citealt{Bergeron92}; \\citealt{Steidel95}; \\citealt{Steidel97}; \\citealt{Churchill05}; \\citealt{Zibetti07}; \\citealt{Chen10}; \\citealt{Nestor11}; \\citealt{Rao11}). Recent work has suggested that the {\\MgII} halos are patchy, with $\\sim$50\\% covering fraction for {\\MgII} (\\citealt{Tripp05}; \\citealt{Churchill07}; \\citealt{Kacprzak08}), and it seems this covering fraction might be even smaller at smaller redshifts (less than 40\\% for strong {\\MgII} absorbers at $z \\sim$~0.1; \\citealt{Barton09}).  Strong {\\MgII} absorbers have been found to be good tracers of cold/warm and low-ionization gas.  \\citet{Rao06} suggests that $\\approx$~80\\% of very strong \\MgII\\, absorption traces damped Lyman-$\\alpha$ (DLA) galaxies ($N(\\HI)>10^{20.3}$), but the relation between $N(\\HI)$ and $N(\\MgII)$ does not prevail for the largest $W_r(\\MgII)$. \\citet{Nestor07} indicates that ``while it is likely that a large fraction of ultrastrong MgII absorption (perhaps~$>$~60\\%) have $N(\\HI)>$ 10$^{20.3}$ atoms cm$^{-2}$, most (perhaps~$\\sim$~90\\%) DLAs have $W_r(\\MgII) \\lsim$~3 {\\AA}''. Indeed, $W_r(\\MgII)$ and $N(\\HI)$ do not correlate as much in this regime and ultrastrong {\\MgII} systems can be found in sub--DLA galaxies (10$^{19}<N(\\HI)<$ 10$^{20.3}$ atoms cm$^{-2}$) as well (\\citealt{Rao11}). Moreover, \\citet{Kulkarni10} has found that the {\\MgII} associated with sub-DLAs tends to show larger velocity spreads on average than that associated with DLAs. Together with the finding that sub-DLAs are more metal-rich than DLAs, \\citet{Kulkarni10} suggests that the difference is due to their star formation histories, where galaxies associated with sub-DLAs would tend to be more massive, and suffer faster star formation and gas consumption, which would result in less $N(\\HI)$ and more metal production. If the star formation has occurred recently, we would expect to see kinematic signatures of it in the {\\MgII} absorption profiles, due to outflows, also called superbubbles or superwinds.  Moreover, we would expect evolution of very strong {\\MgII} absorption profiles as the star formation rate in galaxies shows a decline at redshifts lower than $z \\sim 1$.   Indeed, {\\MgII} absorbers have been observed to show evolution with redshift. In a survey of $\\sim$~3700 quasars from an early data release of Sloan Digital Sky Survey (SDSS), \\citet{Nestor05} found that the number of very strong {\\MgII} absorbers ($W_r > 1$ \\AA) was not consistent with the expectations for cosmological non--evolution, showing larger numbers at high redshift ($z \\gsim$~1.2). This trend is most significant as $W_r(\\MgII)$ increases. \\citet{Prochter06} confirmed these results for the incidence of {\\MgII} in a study of 45,023 quasars from the Data Release 3 (DR3) of SDSS. The decline in number of very strong {\\MgII} absorbers with time is consistent with the decline of cosmic SFR (e.g., \\citealt{Madau98}), which is found to have decreased by about one order of magnitude since its peak (e.g.,\\citealt{RosaGonzalez02} and references therein). Since very strong absorbers are evolving away at redshifts coincident with the decrease in SFR, they could be tracing the star formation sites of galaxies, as previously suggested by \\citet{Guillemin97} and \\citet{Churchill99}. Low-mass galaxies present this peak of SFR at even lower redshifts (\\citealt{Kauffmann04}).  In a study of $z <$~0.8 {\\MgII} absorption of SDSS DR3, \\citet{Bouche06} found an anti-correlation between halo mass of galaxies and the equivalent width of {\\MgII} absorption present. Since the very strong {\\MgII} systems may show saturation in most of their profiles, the equivalent width is determined primarily by the velocity dispersion of the absorbing clouds. The required velocity spreads are consistent with a starburst picture for the strongest {\\MgII} systems, but not with structure within individual virialized halos. \\citet{Prochter06} supported the claim that very strong {\\MgII} absorbing structures are related to superwinds rather than accreting gas in galaxy halos, in a study over a larger redshift range ($0.35 < z < 2.7$) and based on the kinematics of $W_r(\\MgII)>1$~{\\AA} absorbers. High resolution studies of individual absorbers confirm that they show particular signatures in their profiles that might be consistent with superwinds (\\citealt{Bond01}; \\citealt{Ellison03}). Moreover, field imaging of a subset of the strongest {\\MgII} absorbers ($W_r(\\MgII)>2.7$~{\\AA}) at low redshift ($0.42 < z < 0.84$) indicates that interactions, pairs, and starburst related phenomena are likely to be present (\\citealt{Nestor07}).  Weak {\\MgII} absorption systems have been found to evolve as well. \\citet{Narayanan05} and \\citet{Narayanan07} found a peak at $z =$~1.2 in the number per unit redshift, $dN/dz$, of {\\MgII} absorbers with $W_r(\\MgII) <$~0.3 \\AA.  Moreover, \\citet{Narayanan08} find a trend in the ratio of {\\FeII} to {\\MgII}: at higher redshift ($z>$~1.2), weak {\\MgII} absorption systems with large values of $W_r(\\FeII\\lambda2383)/W_r(\\MgII)$ are rare. They suggested that this trend could either be caused by an absence of high density, low ionization gas at high-$z$ or the absence of enrichment by Type Ia supernovae in weak {\\MgII} clouds at high-$z$. They suggest that the relatively few weak {\\MgII} absorbers that are observed at high $z$ are from young stellar populations and thus $\\alpha$--enhanced.  This paper investigates the evolution of very strong {\\MgII} absorption systems with redshift, particularly their {\\FeII} to {\\MgII} ratios.  We enlarge the sample of strong {\\MgII} absorption presented by \\citet{Mshar07}, and compare it to the analysis previously carried out in \\citet{Narayanan07} for weak {\\MgII} absorbers. The data and methods, as well as the systems found, are presented in section \\S\\ref{sec:2}. Results involving evolution of the profiles and the {\\FeII}/{\\MgII} ratio are shown in section \\S\\ref{sec:3}. We summarize the results and discuss their implications in section \\S\\ref{sec:4}. ",
        "conclusions": ""
    },
    "1208/1208.3399_arXiv.txt": {
        "abstract": "There often appear coherently oscillating scalar fields in particle physics motivated cosmological scenarios, which may have rich phenomenological consequences. Scalar fields should somehow interact with background thermal bath in order to decay into radiation at an appropriate epoch, but introducing some couplings to the scalar field makes the dynamics complicated. We investigate in detail the dynamics of a coherently oscillating scalar field, which has renormalizable couplings to another field interacting with thermal background. The scalar field dynamics and its resultant abundance are significantly modified by taking account of following effects: (1) thermal correction to the effective potential, (2) dissipation effect on the scalar field in thermal bath, (3) non-perturbative particle production events and (4) formation of non-topological solitons. There appear many time scales depending on the scalar mass, amplitude, couplings and the background temperature, which make the efficiencies of these effects non-trivial. ",
        "introduction": "\\label{sec:introduction}    Scalar fields play important roles in particle physics and cosmology. The standard model Higgs boson, which has been recently discovered, develops a vacuum expectation value (VEV) and spontaneously breaks gauge symmetry. In the supersymmetric (SUSY) standard model (SM), there are many scalar fields as superpartners of SM quarks and leptons. These scalar fields may have large VEVs in the early Universe and their dynamics may have significant effects on the baryon asymmetry of the Universe through the Affleck-Dine mechanism~\\cite{Affleck:1984fy}.  Inflation in the very early Universe is considered to be caused by a scalar field, called inflaton, which slowly rolls down the scalar potential~\\cite{Lyth:2009zz}. After inflation, the inflaton decays into radiation and then hot thermal Universe begins. The inflaton may also be responsible for the primordial density perturbation which seeds the rich structure of the present Universe. Instead, another scalar field, called curvaton, can explain the primordial density perturbation~\\cite{Mollerach:1989hu,Linde:1996gt,Lyth:2001nq,Moroi:2001ct}. The curvaton also decays into radiation and fluctuations in the curvaton sector are converted to the adiabatic perturbation in the radiation.   In these scenarios, the understandings of the dynamics of coherently oscillating scalar fields are essential for deriving their phenomenological consequences. For example, the inflaton or curvaton must dissipate its energy into the radiation in order for a hot thermal Universe to begin well before the big-bang nucleosynthesis. It inevitably requires interactions between the inflaton/curvaton and other fields, which somehow will be thermalized. A conventional picture of this {\\it reheating} process is that the oscillating scalar field decays into lighter particles and the produced particles are eventually thermalized by gauge and/or Yukawa interactions. This simple picture, however, does not necessarily hold in general. Depending on the couplings between the scalar ($\\phi$) and other fields ($\\chi$), many non-trivial issues appear which makes this topic far more complicated.  Let us assume the simplest Yukawa coupling between the scalar $\\phi$ and a fermion $\\chi$, as \\begin{equation} \\mathcal L = \\lambda \\phi \\bar\\chi_{\\rm L} \\chi_{\\rm R} + {\\rm h.c.},  \\label{yukawa} \\end{equation} in order to induce the $\\phi$ decay into radiation, where $\\lambda$ is a coupling constant and $\\chi$ is assumed to have gauge or Yukawa interactions with SM particles.\\footnote{ For instance, it may be an extra vector-like matter charged under the SM gauge group. ({\\it e.g.}, an extra-quark with a color charge.) As a particular case, $\\chi$ itself may be chiral SM fields. In this case, parametrizing a (flat) direction as $\\phi$ which may have a large initial amplitude, one finds the same interaction as Eq.~\\eqref{yukawa}, given by \\begin{align} {\\cal L} = \\sum_{i,j}\\lambda_{ij} \\phi \\bar\\chi_{{\\rm L},i} \\chi_{{\\rm R},j} +{\\rm h.c.}, \\end{align} where a summation over gauge indices which are not broken by the scalar condensation is promised. } First, the $\\chi$ fields obtain the $\\phi$-dependent mass through this coupling. It is soon recognized that if $\\lambda\\phi \\gg m_\\phi$, where $m_\\phi$ is the mass of $\\phi$, the decay process $\\phi \\to \\chi\\bar\\chi$ is kinematically forbidden except for the region near $\\phi\\sim 0$. In this case, the perturbative decay rate at the true vacuum does not apply and we must take into account particle production phenomena occurring around $\\phi\\sim 0$ and subsequent thermalization processes of $\\chi$. Moreover, in general, there always be background radiation, or thermal bath, even before the reheating completes. Then, if $\\chi$ is light enough, it is thermalized and obtains a thermal mass. If the thermal mass is large enough, the $\\phi$ decay into $\\chi$ is again kinematically forbidden even if the $\\phi$ amplitude is very small. However, even if the decay process is kinematically forbidden, $\\phi$ can dissipate its energy through $\\phi$-$\\chi$ scattering processes. Thermal bath also modifies the effective potential of $\\phi$ through thermal effects and hence the dynamics of $\\phi$ may be significantly affected.   Therefore, even only the introduction of a simple interaction (\\ref{yukawa}) leads to many complicated issues on the dynamics of $\\phi$. Without careful treatment of these issues, we cannot discuss cosmological effects of $\\phi$. To the best of our knowledge, however, there have not been comprehensive and complete analyses on this topic. It was only partly attacked by separate literature as follows.  \\begin{itemize}  \\item Effective potential of the scalar field in thermal environment is well known in the case where $\\chi$ takes part in thermal bath~\\cite{KapustaBook}. On the other hand, thermal modification on the effective potential when $\\chi$ decouples from thermal bath at large $\\phi$ value seems to be less known except for the context of Affleck-Dine mechanism~\\cite{Anisimov:2000wx}. However, this modification is rather a generic feature if $\\chi$ has a gauge/Yukawa interaction and needs to be included in the analyses.  \\item Besides thermal modification on the effective potential, $\\phi$ also receives thermal dissipation. Dissipation of oscillating scalar field in thermal bath was extensively studied in the context of warm inflation in Refs.~\\cite{Berera:1995ie,Berera:2008ar,BasteroGil:2009ec}, and also in Refs.~\\cite{Yokoyama:2004pf,Drewes:2010pf} in a context of inflation. Ref.~\\cite{BasteroGil:2010pb} also considered the dissipation of scalar fields in a setup close to ours. However, the dynamical aspects of oscillating scalar field with a large amplitude were not considered.  \\item The effect of non-perturbative particle production, or called {\\it preheating}, was studied in detail in Refs.~\\cite{Kofman:1994rk,Shtanov:1994ce,Kofman:1997yn} and thereafter. Most of these studies focused on the $\\chi$ particle production neglecting the $\\chi$ interaction with thermal bath. Ref.~\\cite{Felder:1998vq} considered the case where produced $\\chi$ particles decay into radiation: so-called {\\it instant preheating}. This may efficiently transfers the $\\phi$ energy density to radiation even in the case where $\\phi$ has a large amplitude and the standard calculation of perturbative decay into $\\chi$ is not applicable.  \\item In association with thermal modification on the effective potential, the scalar field may fragment into the non-topological solitons. Once most energy of the oscillating scalar field is absorbed by solitons, estimation of the dissipation rate and the decay temperature of $\\phi$ may be significantly modified. This is studied in the context of Q-balls when the complex scalar oscillates with an elliptical orbit in the complex plane. The formation of solitons, however, can occur even in the case of a real scalar. Its phenomenological consequences have not been investigated in detail in a setup we are interested in.  \\end{itemize}  In this paper, we consider the dynamics of $\\phi$ having interaction of the type (\\ref{yukawa}) in thermal environment. We take into account various effects listed above in reliable manners and study in which situation these effects become substantial for the scalar dynamics. As for the issue of non-topological solitons, we avoid to give definite conclusions but mention possible consequences. We only have four model parameters: $m_\\phi, \\phi_i, \\lambda, \\alpha (\\equiv g^2/4\\pi)$, where $m_\\phi$ is the bare mass of $\\phi$, $\\phi_i$ is the initial amplitude of $\\phi$, $\\lambda$ is the Yukawa coupling between $\\phi$ and $\\chi$, $g$ is the coupling constant between $\\chi$ and thermal bath, which is typically an SM gauge coupling. In addition, we introduce the reheating temperature $T_{\\rm R}$, which controls the density of thermal bath.\\footnote{ In this paper we assume the existence of thermal bath, which comes from the inflaton decay. Thus our analysis of the scalar dynamics is not applied to the inflaton itself. The inflaton dynamics is also an interesting subject, which will be discussed elsewhere. \\label{foot:inflaton} } Even in this simple setup, there appear many time scales and we must be cautious about the usage of a particular formalism case-by-case. Eventually, we have successfully figured out the global picture of the scalar dynamics in broad parameter spaces. We believe that, although this is a simplified toy model, it captures essential features of the most applications to realistic models. This is because oscillating scalar fields must decay into radiation unless it becomes (a part of) dark matter, and then it inevitably couples to lighter species, which in turn have interactions with SM particles, and this is nothing but a situation we are considering. Even if $\\phi$ couples to many fields with coupling strength varying by orders of magnitude, the dynamics mainly depends on the largest coupling and hence our analyses apply.  This paper is organized as follows. In Sec.~\\ref{sec:over} we describe our basic setup and list some particle physics motivated examples. In Sec.~\\ref{sec:thermal} we review thermal effects on the scalar field evolution based on the closed time path formalism. In particular, thermal modification on the effective potential and the dissipation rate will be introduced, which significantly affects the scalar dynamics. In Sec.~\\ref{sec:non}, the effects of non-perturbative particle production are discussed. This also serves as a dissipation of the oscillation of the scalar field. In Sec.~\\ref{sec:dyn}, we solve the scalar dynamics taking these effects into account. Readers who are only interested in results may skip preceding sections and only read Sec.~\\ref{sec:dyn}. We conclude in Sec.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}    In this paper we have extensively studied the scalar dynamics in the early Universe taking account of the effects of thermal environment. Despite the fact that the early Universe after inflation may be filled with high-temperature thermal plasma, thorough analyses on the scalar dynamics in such an environment were missing in the literature. Needless to say, scalar fields play important roles in cosmology. Inflaton explains the primordial inflation and the primordial density perturbation, and curvaton may also be responsible for the generation of density perturbation. Affleck-Dine fields in the MSSM can create the baryon asymmetry of the Universe. SUSY breaking fields may have significant effects on cosmology. Similar long-lived scalars often appear in extensions of the SM. All these scalar fields generally have large energy densities so that they must decay or evaporate at an appropriate epoch in order not to disturb the success of standard cosmology. Therefore, they necessarily couple to SM particles directly or indirectly through some intermediate states.  Based on these observations, we considered a model in which a scalar $\\phi$ couples to fermions $\\chi$ through a Yukawa coupling, which then interact with thermal bath. This simple model captures essential features of realistic models. We have consistently taken into account various effects: thermal modification on the effective potential of $\\phi$, thermal dissipation of $\\phi$, non-perturbative particle production and the formation of non-topological solitons. It is found that even in this simple class of models, the scalar dynamics is so complicated that a naive estimation neglecting these effects is not allowed in broad parameter spaces. In particular, it should be noticed that thermal dissipation, which can be interpreted as the $\\phi$ evaporation due to scatterings with particles in an environment, often plays a dominant role in determining the dissipation epoch of $\\phi$.  Finally, let us mention what we have not included in our analyses. In this paper we have followed the evolution of only the zero-mode of $\\phi$ coherent oscillation. However, it is expected that fluctuations around the zero-mode, or the particle-like excitations of $\\phi$ can be important if thermal effects play a dominant role for the evaporation of $\\phi$. This can in principle be traced by following the evolution of the two-point functions of $\\phi$. We have also restricted ourselves to the case of $\\lambda < \\alpha$, because otherwise the whole dynamics including $\\chi$ particle production is complicated. We will return to these issues in some concrete models elsewhere~\\cite{mukaida}."
    },
    "1208/1208.0679_arXiv.txt": {
        "abstract": "We present first results from follow-up of targets in the northern hemisphere $\\beta$ Pictoris and AB Doradus moving group candidate list of Schlieder, L\\'epine, and Simon~(2012).  We obtained high-resolution, near-infrared spectra of 27 candidate members to measure their radial velocities and confirm consistent group kinematics.  We identify 15 candidates with consistent predicted and measured radial velocities, perform analyses of their 6-dimensional (\\emph{UVWXYZ}) Galactic kinematics, and compare to known group member distributions.  Based on these analyses, we propose that 7 $\\beta$ Pic and 8 AB Dor candidates are likely new group members.  Four of the likely new $\\beta$ Pic stars are binaries; one a double lined spectroscopic system.  Three of the proposed AB Dor stars are binaries.  Counting all binary components, we propose 22 likely members of these young, moving groups.  The majority of the proposed members are M2 to M5 dwarfs, the earliest being of type K2.  We also present preliminary parameters for the two new spectroscopic binaries identified in the data, the proposed $\\beta$ Pic member and a rejected $\\beta$ Pic candidate.  Our candidate selection and follow-up has thus far identified more than 40 low-mass, likely members of these two moving groups.  These stars provide a new sample of nearby, young targets for studies of local star formation, disks and exoplanets via direct imaging, and astrophysics in the low-mass regime. ",
        "introduction": "Young stars are critical for understanding star formation and evolution and the circumstellar environment.  For example, the $\\sim$10 Myr old A type dwarf $\\beta$ Pictoris is host to the first directly imaged debris disk (Smith and Terrile~1984).  The first direct images of extrasolar planets were also achieved around nearby, young stars.  Images of the planetary system around the $\\sim$30 Myr A type dwarf HR 8799 have provided empirical evidence to constrain planetary formation models and the evolution of the atmospheres in young, substellar objects (Marois et al.~2008; Marois et al.~2010; Barman et al.~2011, Skemer et al.~2012).  $\\beta$ Pictoris also hosts an imaged massive planet whose orbit is consistent with an observed warp in the disk (Lagrange et al.~2010, 2012).  The previous stellar examples are members of the youngest populations of stars in the solar neighborhood; nearby, young moving groups (\\emph{NYMGs}, Zuckerman and Song~2004, hereafter ZS04; Torres et al.~2008, hereafter T08).  Although moving group members have aged enough ($\\sim$10-100 Myr) to be dispersed from their formation region, they retain small velocity dispersions and their common galactic motion translates to proper motions toward a convergent point in the plane of the sky (see Jones~1971; de Bruijne~1999; Mamajek~2005).  Because of their proximity (members $\\lesssim$100 pc) and loosely bound nature, moving groups can be spread over thousands of square degrees on the sky.  Known groups include the AB Doradus and $\\beta$ Pictoris moving groups ($\\sim$50-150 and 10-20 Myrs old respectively), which are the only groups to have substantial numbers of known members in the north.  Tracing back the kinematics of the youngest \\emph{NYMGs}, including $\\beta$ Pic, shows that they may be related to a common star formation event in the Sco-Cen region (Mamajek, Lawson, and Feigelson~2000; Mamajek and Feigelson~2001, Fern\\`andez et al.~2008, hereafter F08). In contrast, the older AB Dor group has kinematics and age similar to the Pleiades.  Luhman, Stauffer, and Mamajek (2005) suggest that AB Dor and the Pleiades are linked to the same star formation event.  Member populations in these groups remain largely incomplete and searches to find the missing members are ongoing (e.g.~L\\'epine and Simon~2009, hereafter LS09; Looper et al.~2010a, 2010b; Rice, Faherty, and Cruz~2010; Schlieder, L\\'epine, and Simon~2010, hereafter SLS10; Kiss et al.~2011;  Riedel et al.~2011; Rodriguez et al.~2011; Shkolnik et al.~2011; Zuckerman et al.~2011; Malo et al.~2012; Schlieder, L\\'epine, and Simon~2012, hereafter SLS12; Shkolnik et al.~2012).  Low-mass members (M$\\lesssim$1.0 M$_{\\odot}$) remain the most under sampled population because: (1) the intrinsic luminosity of low-mass stars places them beyond the flux limits of the X-ray and astrometric catalogs that have served as the primary resources for identifying new group members (see ZS04; T08 and references therein), and (2) Youth diagnostics for the lowest-mass stars lack reliability, this makes the missing stars more difficult to identify from the field (see Schlieder et al.~2012, hereafter S12).  The low-mass populations of moving groups are key for studies of recent local star formation, disks and exoplanets, and the astrophysics of young, low-mass stars.  With a more complete moving group census the mass functions of individual groups could be determined, allowing direct comparison to young open clusters and the field (e.g. Kroupa et al.~2002, Chabrier et al.~2005, Bochanski et al.~2010).  Studies of this kind could provide insight into the universality of the initial mass function and reveal information about the formation environment of \\emph{NYMGs}.   Since moving groups may represent the endpoint of clustered star formation, they are crucial not only for understanding the recent star formation history near the Sun, but the latter stages of cluster evolution and dispersion in general.  A complete sample of nearby, young stars would also allow for detailed multiplicity studies.   Multiplicity fractions, mass-ratios, and semi-major axis distributions can also be linked to formation environment and dynamical evolution (see Chauvin et al.~2010; Bergfors et al.~2010; Janson et al.~2012, hereafter J12, and references therein).  In addition, identification of low-mass moving group members will allow for the detailed investigation of signatures of youth in low-mass stars.  These include chromospheric and coronal activity indicators, lithium abundances, gravity sensitive spectral features, and rapid rotation.  A more complete sample would be particularly useful in understanding the evolution and reliability of these youth signatures at different ages across the transition to the fully convective regime ($\\sim$M4 spectral type, see West et al.~2011, hereafter W11; Reiners and Mohanty~2012, hereafter RM12; Reiners, Nandan, and Goldman~2012; S12).  Observations of the few known, low-mass \\emph{NYMG} members have provided a wealth of information regarding circumstellar disks and substellar companions.  The most classic example may be AU Mic, a $\\sim$10 Myr old M dwarf in the $\\beta$ Pic group, which hosts an edge on debris disk that has acted as a benchmark for disk evolution in the low-mass regime (Kalas et al.~2004, Liu~2004).  Known low-mass members host to substellar companions include 2MASS 1207 in the TW Hydrae association, an $\\sim$8 Myr old brown dwarf with a planetary mass companion (Chauvin et al.~2004), and CD-35 2722 and 1RXS J2351+3127, AB Dor M dwarfs with brown dwarf companions (Wahhaj et al.~2011; Bowler et al.~2012, hereafter B12).  Further imaging surveys of low-mass moving group stars will provide statistics on the fraction of M dwarfs with massive planet and brown dwarf companions (see B12; Delorme et al.~2012).  These observations can infer key constraints on multiple system and planet formation in the low-mass regime.  A new technique to identify low-mass moving group members was described in LS09 and SLS10.  Low-mass candidate group members are selected on the basis of proper motion, photometry, and activity.  In SLS12 we described a comprehensive list of young $\\beta$ Pictoris and AB Dor candidates in the northern hemisphere selected using this technique.  The use of new proper motion catalogs and UV flux as an activity indicator allowed access to large numbers of previously inaccessible, young candidates.  These stars require detailed follow-up to verify group membership.  In this paper we present our first results from a multi-epoch, multi-telescope follow-up campaign of this candidate list.  In \\S 2 we describe our observations and present the results.   An analysis of the 6-dimensional (6D, \\emph{UVWXYZ}) kinematics of proposed likely members is presented in \\S 3.  Likely new members (\\emph{LNMs}) are treated individually in \\S 4 and two new spectroscopic binaries from the sample are discussed in more detail in \\S 5.  We present a discussion of proposed member distances and ages in \\S 6 and a summary follows in \\S 7. ",
        "conclusions": "\\subsection{Proposed Member Distances}  It is important to reiterate that our kinematic analyses used the $d_{kin}$ for all candidates.  This distance is an estimate assuming group membership.  We have verified that these distances are consistent with group cluster sequences in color-magnitude diagrams (see SLS12) and consistent with models at the ages of the groups.  For most proposed new members, photometric or spectroscopic distance estimates are available in the literature, but have varying levels of uncertainty (see \\S\\ref{NORTH_LNM}).  Since the (\\emph{UVWXYZ}) calculations are directly dependent on stellar distance, trigonometric parallax measurements from ongoing programs\\footnote{Follow-up parallax observations are in progress at the MDM observatory (S. L\\'epine, private communication) and US Naval Observatory (H. Harris, private communication).} will determine whether or not \\emph{LNM} kinematics are truly consistent with group membership.  \\subsection{Proposed Member Ages}  In SLS12 we defined two types of candidate based on their (\\emph{V-K$_{s}$}) colors.  We designated candidates with (\\emph{V-K$_{s}$})$\\le$5.0 candidates with reliable youth (\\emph{CWRYs}), and candidates with (\\emph{V-K$_{s}$})$>$5.0, candidates with ambiguous youth (\\emph{CWAYs}).  These designation were designed to reflect the observed and modeled steep transition in M dwarf activity (H$\\alpha$) lifetimes at approximately M4 SpTy (see W11, RM12).  Stars earlier than M4 are active only when young and stars at later types have increasing activity lifetimes, up to $\\sim$8 Gyr for the latest M's.  Thus, the activity used as the primary indicator of youth in the SLS12 list has varying levels of reliability depending on candidate spectral type.  Here we discuss the implications of this and other age diagnostics on our proposed members.  Our activity based youth selection is modeled after the methods of Shkolnik et al.~(2009, 2011).   They estimate that X-ray and UV activity consistent with other, young late-type stars places an upper limit of $\\sim$300 Myr on the age of an M dwarf. This limit is more reliable for the stars that qualify as \\emph{CWRYs}; all but four of the proposed new members.   We can place an age constraint of $<$100 Myr on two of the \\emph{CWAY} new members,  PYC J05019+0108 and PYC J21376+0137, based on weak Na 8200~\\AA~doublets (see S12).  Other searches for new \\emph{NYMG} members have used the EW of the Li feature at 6708~\\AA~as a requirement for membership.  For example, Shkolnik et al.~(2011, 2012) used Li EW as cut in their search for new, low-mass members of moving groups.  They found fewer new, M dwarf members than anticipated and proposed that Li depletion may be too strict a requirement for membership.  Li depletion, like activity, has been shown both empirically and theoretically to be unreliable when estimating low-mass star ages (see Yee and Jensen~2010, Baraffe and Chabrier~2010).  Furthermore, for the case of the older AB Doradus moving group, most M dwarfs will surely have depleted their lithium (see Mentuch et al.~2008)  Thus, the stars we propose as likely new members of the $\\beta$ Pic and AB Dor groups are pre-main sequence, but whether their ages meet the strict criteria imposed by higher mass group members is difficult to determine at this time.  Ongoing spectroscopic follow-up to investigate Li depletion and gravity sensitive spectral features must be combined with existing measurements of photometry, activity, and rotation to build an overarching view of youth on a star by star basis.  The stars we are trying to identify as new moving group members are exactly those stars that elude accurate and reliable age dating. Therefore, the nearby, young M dwarfs presented here are critical to lay the groundwork for new, more reliable methods of age dating low-mass stars."
    },
    "1208/1208.2951_arXiv.txt": {
        "abstract": "We calculate the rate for thermal production of axions and saxions via scattering of quarks, gluons, squarks, and gluinos in the primordial supersymmetric plasma. Systematic field theoretical methods such as hard thermal loop resummation are applied to obtain a finite result in a gauge-invariant way that is consistent to leading order in the strong gauge coupling. We calculate the thermally produced yield and the decoupling temperature for both axions and saxions. For the generic case in which saxion decays into axions are possible, the emitted axions can constitute extra radiation already prior to big bang nucleosynthesis and well thereafter. We update associated limits imposed by recent studies of the primordial helium-4 abundance and by precision cosmology of the cosmic microwave background and large scale structure. We show that the trend towards extra radiation seen in those studies can be explained by late decays of thermal saxions into axions and that upcoming Planck results will probe supersymmetric axion models with unprecedented sensitivity. ",
        "introduction": "There are several hints towards physics beyond the standard model (\\SM). One of them is the strong CP problem. If this problem is solved via the Peccei--Quinn (\\PQ) mechanism, the axion $\\axion$ arises as the pseudo-Nambu-Goldstone boson associated with the U(1)$_\\mathrm{PQ}$ symmetry broken spontaneously at the \\PQ\\ scale $\\fax$~\\cite{Sikivie:2006ni,Kim:2008hd}. Another attractive extension of the \\SM\\ is supersymmetry (SUSY)~\\cite{Martin:1997ns,Drees:2004jm,Baer:2006rs,Dreiner:2008tw}. In conceivable settings with both the \\PQ\\ mechanism and SUSY, the pseudo-scalar axion is part of a supermultiplet in which also its scalar partner, the saxion $\\saxion$, and its fermionic partner, the axino $\\axino$, appear. The energy density of the early Universe can then receive contributions from coherent oscillations of the axion field~\\cite{Beltran:2006sq,Sikivie:2006ni,Kim:2008hd} and the saxion field~\\cite{Chang:1996ih,Hashimoto:1998ua,Asaka:1998ns,Kawasaki:2007mk,Kawasaki:2011aa} and from thermal production of axions~\\cite{Turner:1986tb,Chang:1993gm,Masso:2002np,Hannestad:2005df,Sikivie:2006ni,Graf:2010tv}, saxions~\\cite{Kim:1992eu,Chang:1996ih,Asaka:1998ns}, and axinos~\\cite{Rajagopal:1990yx,Bonometto:1993fx,Chun:1995hc,Asaka:2000ew,Covi:2001nw,Brandenburg:2004du,Strumia:2010aa,Chun:2011zd,Bae:2011jb,Choi:2011yf,Bae:2011iw} in the hot primordial plasma.  Here we calculate for the first time the thermal production rate of axions and saxions via scattering processes of quarks, gluons, squarks, and gluinos in a gauge-invariant way consistent to leading order in the strong coupling constant $g_s$. In our calculation we use hard thermal loop (HTL) resummation~\\cite{Braaten:1989mz} and the Braaten--Yuan prescription~\\cite{Braaten:1991dd} to account systematically for screening effects in the quark-gluon-squark-gluino plasma (QGSGP). This method was introduced on the example of axion production in a hot QED plasma~\\cite{Braaten:1991dd}; see also~Ref.~\\cite{Bolz:2000fu}. Moreover, it has been applied to calculate the thermal production of gravitinos~\\cite{Bolz:2000fu,Pradler:2006qh,Pradler:2006hh,Pradler:thesis} and axinos~\\cite{Brandenburg:2004du} in SUSY settings and of axions in a non-SUSY quark-gluon plasma (QGP)~\\cite{Graf:thesis,Graf:2010tv}.  Based on our result for the thermal axion/saxion production rate, we determine the respective thermally produced yields and estimate the decoupling temperature of axions and saxions from the thermal bath. While both axions and axinos are promising dark matter candidates (cf.~\\cite{Sikivie:2006ni,Kim:2008hd,Steffen:2008qp,Covi:2009pq} and references therein), saxions can be late decaying particles with potentially severe cosmological implications. For example, energetic hadrons and photons from saxion decays during or after big bang nucleosynthesis (BBN) can change the abundances of the primordial light elements~\\cite{Kawasaki:2007mk}. Moreover, photons from saxion decays can affect the black body spectrum of the cosmic microwave background (CMB) for a saxion lifetime of $10^6\\,\\seconds\\lesssim \\tausaxion\\lesssim 10^{13}\\,\\seconds$~\\cite{Asaka:1998xa,Kawasaki:2007mk} or may contribute either to the diffuse $X(\\gamma)$-ray background or as an additional source of reionization for $\\tausaxion\\gtrsim 10^{13}\\,\\seconds$~\\cite{Kawasaki:1997ah,Asaka:1998xa,Chen:2003gz,Kawasaki:2007mk}. In scenarios in which the decay mode into axions is not the dominant one, saxion decays may also produce significant amounts of entropy~\\cite{Kim:1992eu,Lyth:1993zw,Hashimoto:1998ua, Asaka:1998xa,Hasenkamp:2010if,Baer:2010gr}. This can dilute relic densities of species decoupled from the plasma and also the baryon asymmetry $\\eta$. Then, $\\tausaxion<1~\\seconds$ is imposed by successful BBN which requires a standard thermal history for temperatures below $T\\sim1~\\MeV$.  In this work, however, we look at scenarios in which saxions (from thermal processes) decay predominantly into axions. Moreover, we still focus on decays prior to BBN and compute the additional radiation provided in the form of the emitted relativistic axions. Such a non-standard contribution $\\Delta\\neff$ to the effective number of light neutrino species $\\neff$ from decays of thermal saxions into axions was previously considered in Refs.~\\cite{Chun:1995hc,Chang:1996ih,Choi:1996vz,Kawasaki:2007mk}. Applying our new result for the thermally produced saxion yield and new cosmological constraints on $\\Delta\\neff$ imposed by recent studies of BBN, the CMB, and large scale structure (LSS)~\\cite{Izotov:2010ca,Aver:2010wq,Hamann:2010pw}, we present updated limits on the \\PQ\\ scale $\\fax$, the saxion mass $\\msaxion$, and the reheating temperature $\\TR$ after inflation.  Interestingly, precision cosmology~\\cite{Hamann:2007pi,Reid:2009nq,Komatsu:2010fb,Hamann:2010pw,GonzalezGarcia:2010un} and recent studies of the primordial $^4$He abundance~\\cite{Izotov:2010ca,Aver:2010wq} show a trend towards a radiation content that exceeds the predictions of the SM. In fact, such an excess can be explained by the considered saxion decays into axions. The observed trend may thus be a hint for the existence of a SUSY axion model. Here results from the Planck satellite mission will be extremely valuable, which will come with an unprecedented sensitivity to the amount of extra radiation at times much later than those at which BBN probes this quantity. Based on a forecasted 68\\% confidence level (CL) sensitivity of $\\Delta\\neff=0.26$~\\cite{Perotto:2006rj,Hamann:2007sb}, we indicate parameter regions of SUSY axion models that will be tested by results from the Planck satellite mission expected to be published in the near future.  The remainder of this paper is organized as follows. In the next section we consider interactions of the PQ supermultiplet and decay widths for saxion decays. In Sects.~\\ref{Sec:ThermalSaxionProduction} and~\\ref{Sec:ThermalAxionProduction} our calculations of the thermal production rates of saxions and axions are presented. We compute the associated yields in Sect.~\\ref{Eq:SaxionAxionYield} and use the results to estimate the saxion/axion decoupling temperature in Sect.~\\ref{Sec:TD}. Then we explore $\\Delta\\neff$ provided in the form of axions from saxion decays and possible manifestations in studies of BBN and of the CMB and LSS. Here we comment on potential restrictions which can emerge from overly efficient thermal gravitino/axino production and describe exemplary settings that allow for a high reheating temperature of $\\TR\\sim 10^8$--$10^{10}\\,\\GeV$. In Sect.~\\ref{Sec:AxionDensity} we compare the relic density of axions from the misalignment mechanism with the ones of thermal axions and of non-thermal axions from saxion decays. Our conclusions are given in Sect.~\\ref{Sec:Conclusion}. ",
        "conclusions": "\\label{Sec:Conclusion}  We have explored thermal production processes of axions and saxions in the primordial plasma, resulting axion populations and their manifestations in the form of extra radiation $\\Delta\\neff$ prior to BBN and well thereafter. The considered SUSY axion models are attractive for a number of reasons. For example, they allow for simultaneous solutions of the strong CP problem, the hierarchy problem, and the dark matter problem.  Here we have focussed on the saxion, which can be a late decaying particle and as such be subject to various cosmological constraints. We find that the saxion decay into two axions is often the dominating one. For a saxion mass of $\\msaxion\\gtrsim 1~\\GeV$, such decays occur typically before the onset of BBN. We have shown that the emitted axions can then still be relativistic at the large scattering surface. Thereby, they can provide sizable contributions to extra radiation $\\Delta\\neff$ that is testable in BBN studies and in precision cosmology of the CMB and the LSS.  We have aimed at a consistent description of both the thermal axion/saxion production and of saxion decays into axions. This has motivated our careful derivations of the Lagrangian $\\mathcal{L}_\\text{\\PQ}^\\text{int}$ that describes the interactions of the PQ multiplet with quarks, gluons, squarks, and gluinos and of $\\mathcal{L}_\\text{\\PQ}^\\text{kin}$ that describes the interactions of saxions with axions in addition to their kinetic terms. The requirement of canonically normalized kinetic terms defines the scale $\\vax$, which governs the saxion-axion-interaction strength together with another PQ-model-dependent parameter $x\\lesssim 1$. On the other hand, the form of the effective axion-gluon-interaction term defines the PQ scale $\\fax$. Considering the emergence of this term from loops of heavy KSVZ fields in an explicit hadronic axion model, we find $\\fax=\\sqrt{2}\\vax$. This is in contrast to numerous existing studies which treat $\\vax$ and $\\fax$ synonymously.  Relying on the derived form of $\\mathcal{L}_\\text{\\PQ}^\\text{int}$, we have calculated the thermal production rates of saxions and axions and the resulting yields in the hot early Universe. Despite differences in the interaction terms, we find that the rate for thermal saxion production agrees with the one for thermal axion production. This implies an agreement also of the calculated thermally produced yields and of our estimates of the decoupling temperatures $\\TD$. By applying HTL resummation~\\cite{Braaten:1989mz} and the Braaten--Yuan prescription~\\cite{Braaten:1991dd}, finite results are obtained in a gauge-invariant way consistent to leading order in the coupling constant % and screening effects are treated systematically.  Using our result for the thermally produced saxion yield, we have calculated $\\Delta\\neff$ provided in the form of axions from decays of thermal saxions. This has allowed us to demonstrate that such a  $\\Delta\\neff$ contribution can indeed explain the trends towards extra radiation beyond the SM seen in recent studies of BBN, CMB, and LSS.  To account for the current PDG recommendation for the free neutron lifetime, $\\tau_{\\mathrm{n}}=880.1\\pm1.1~\\seconds$~\\cite{Beringer}, we have performed a BBN likelihood analysis with {\\tt PArthENoPE}~\\cite{Pisanti:2007hk} and based on recent insights on the primordial abundances of $^4$He~\\cite{Izotov:2010ca,Aver:2010wq} and D~\\cite{MNR:MNR13921}. For $\\Delta\\neff$ at the onset of BBN, we thereby obtain posterior maxima of 0.76 and 0.77 and $3\\sigma$ upper limits of 1.97 and 3.53 with the $Y_\\text{p}$ results of~\\cite{Izotov:2010ca} and~\\cite{Aver:2010wq}, respectively. When comparing these values with results from studies of the CMB and LSS, we find that the latter provide compatible but more pronounced hints for extra radiation. For example, the precision cosmology study of~\\cite{Hamann:2010pw} reports a mean of 1.73 and a $2\\sigma$ limit 3.59 for $\\Delta\\neff$ at $T\\ll 1~\\MeV$ when using the CMB + HPS + HST data set.  We have translated the upper limits on $\\Delta\\neff$ quoted above into bounds on $\\fax$, $\\msaxion$, and $\\TR$. These bounds can disfavor significant regions of the $\\msaxion$--$\\fax$ parameter plane in high $\\TR$ scenarios. However, we find that our limits leave open a considerable parameter region previously thought to be excluded~\\cite{Kawasaki:2007mk}. Significant parts of the allowed parameter region have been identified, which will become accessible very soon with the upcoming results from the Planck satellite mission.  The explanation of the above hints for extra radiation via axions from decays of thermal saxions requires a relatively high reheating temperature of $\\TR\\gtrsim 10^7~\\GeV$ for $\\msaxion\\gtrsim 0.1~\\GeV$. Such high $\\TR$ scenarios can be in conflict with cosmological constraints due to overly efficient thermal production of axinos and gravitinos. To illustrate the viability of $\\Delta\\neff\\sim 1$ from saxion decays, we have described two exemplary SUSY scenarios which allow for $\\TR=10^8$ and $10^{10}~\\GeV$: \\begin{itemize} \\item[(i)] With the gravitino LSP as cold dark matter and a sneutrino NLSP, the presented $\\Delta\\neff\\sim 1$ explanation for $\\TR=10^8~\\GeV$ can be viable for $\\msaxion\\sim\\mgravitino\\sim 10~\\GeV$ and $m_{\\tilde{g}}\\sim 1~\\TeV$. This explanation requires $\\fax\\sim 10^{10}\\,\\GeV$ and heavy axinos, $\\maxino\\gtrsim 2~\\TeV$, which decay prior to NLSP decoupling. Here $\\Delta\\neff\\sim 1$ is already present at the onset of BBN and does not change thereafter. Accordingly, we expect that the Planck results will point to a $\\Delta\\neff$ value that is consistent with the one inferred from BBN studies. \\item[(ii)] With a very light axino LSP, $\\maxino\\lesssim 37~\\eV$, as hot dark matter and a gravitino NLSP, the $\\Delta\\neff\\sim 1$ explanation for $\\TR=10^{10}~\\GeV$ can be viable for $\\msaxion\\sim\\mgravitino\\sim 100~\\GeV$ and $m_{\\tilde{g}}\\sim 1~\\TeV$. Here this explanation requires $\\fax\\sim 10^{12}\\,\\GeV$ so that cold dark matter can be provided by the axion misalignment mechanism. With the stau as the LOSP, further potential BBN constraints can be evaded. Note that $\\TR=10^{10}~\\GeV$ allows for successful thermal leptogenesis. The saxion decays give $\\Delta\\neff\\sim 1$ already at the onset of BBN. However, late gravitino NLSP decays into the axion and the axino LSP can provide an additional contribution of $\\Delta\\neff\\sim 1$ well after BBN~\\cite{Ichikawa:2007jv,Hasenkamp:2011em}. Thus, it will be interesting to see whether the Planck results confirm the trend towards an excess of extra radiation that is more pronounced at late times. For example, the finding of $\\Delta\\neff\\sim 2$ at late times will be a possible signature expected in this setting. \\end{itemize}  If a SUSY hadronic axion model is realized in nature, three different axion populations will be present today: thermally produced/thermal relic axions, non-thermally produced axions from decays of thermal saxions, and the axion condensate from the misalignment mechanism. We have calculated and compared the associated density parameters. The results allow us to infer the axion analog of the Lee--Weinberg curve. For $\\fax\\gtrsim 10^{12}~\\GeV$ and an initial misalignment angle of $\\theta_i\\sim 1$, the axion density parameter is governed by the axion condensate. In that parameter region this population may be accessible in direct axion dark matter searches. For smaller $\\fax$ and smaller $\\theta_i$, axions from saxion decays can dominate the axion density parameter. While it will be extremely challenging to probe thermal axions, Planck may confirm $\\Delta\\neff$ signals of this non-thermally produced population in the full allowed $\\fax$ range. Since the considered axion populations can coexist, there is the exciting chance to see signals of both axion dark matter and axion dark radiation in current and future experiments."
    },
    "1208/1208.2160_arXiv.txt": {
        "abstract": "{} {We determine the number density and area contribution of small-scale inter-granular Ca\\,{\\small II} bright G-band structures in images of the quiet Sun as tracers of kilo-Gauss magnetic flux-concentrations.} {In a $149\\arcsec\\,\\times\\,117\\arcsec$ G-band image of the disk center at the activity minimum, 7593 small inter-granular structures were segmented with the `multiple-level tracking' pattern recognition algorithm. The scatterplot of the continuum versus the G-band brightness shows the known magnetic and non-magnetic branches. These branches are largely disentangled by applying an intrinsic Ca\\,{\\small II}\\,H excess criterion. The thus obtained 2995 structures contain 1152 G-band bright points (BP) and 1843 G-band faint points (FP). They show a tendency toward increasing size with decreasing G-band excess, as expected from the `hot wall' picture. Their Ca\\,{\\small II}\\,H and G-band brightness are slightly related, resembling the known relation of Ca\\,{\\small II} and magnetic field strength. The magnetic flux density of each individual BP and FP is estimated from their G-band brightness according to MHD model calculations.} {The entity of BP and FP covers the total FOV with a number density of 0.32\\,/\\,Mm$^2$ and a total area contribution of 2.0\\%. Their individual calibrations yield a mean flux density of 20\\,Mx/cm$^2$ in the entire FOV and 13\\,Mx/cm$^2$ for inter-network regions.} {} ",
        "introduction": "The small-scale magnetic flux-concentrations on the non-active solar surface are an important ingredient for understanding solar magnetism. Active region features, such as sunspots, faculae, prominences, and flares, exhibit a strong variation with the solar cycle, commonly explained by a global solar dynamo acting in the deep convection zone. On the other hand, small-scale magnetic flux-concentrations exist that cover the solar surface more homogeneously in time and space indicating a less tight dependence on the global dynamo  Unfortunately, it is difficult to measure magnetic flux on those scales predicted by MHD modeling (V\\\"ogler \\& Sch\\\"ussler 2007). Measurements of circular polarization due to the Zeeman effect can only give the {\\it net} flux, since opposite polarities cancel each other out within the spatial resolution element (Dom\\'inguez Cerde\\~na et al. 2006). The absolute (unsigned) flux, however, has to be deduced from magnetically broadened intensity profiles (Stokes-I) that are affected not only by the magnetic field but also by the, a priori unknown, intrinsic atmosphere of each individual flux concentration.  The problems for such a conversion have been discussed by, e.g., S\\'anchez Almeida (2003). Algorithms used for this procedure are often called `inversion codes', which might suggest that the observed polarization profiles are actually `inverted' into magnetic field parameters. Instead, these codes vary both the magnetic field structure and the intrinsic atmosphere model until an optimal parameter fit between calculated and observed profiles is found. The codes also optimize the filling factor of the magnetic signal inside the resolution element but do not consider two-dimensional radiative effects in very small structures.  Recent observations impressively show that flux-concentrations can be of such a small size that even the 0.09\\arcsec{} spatial resolution of the 1\\,m SST at La Palma (Scharmer et al. 2003) does not reach the lower end of the size distribution of G-band structures. Stokes images, however, cannot achieve such a high spatial resolution, due to additional optical elements and longer exposure times. Hence, they even depend more on the (doubtful) fill factor than intensity images.  Many of the problems in interpreting polarization measurements may be avoided by analyzing intensity images (Sanchez Almeida et al. 2004; De Wijn et al. 2005). Small-scale magnetic concentrations of kilo-Gauss flux density are known to exhibit excess brightness in the CH-band near 430\\,nm (G-band) over the neighboring continuum (Berger et al. 1995), thus allowing their study in post factum reconstructed images of higher spatial and temporal resolution than achieved for Stokes measurements. The excess in the G-band brightness over the neighboring spectral continuum is a different measure than the excess over the mean photospheric G-band level, defining G-band bright points (BP; Muller \\& Roudier 1984). The latter cover two families of structures: (I) inter-granular BP, which are magnetic (Berger \\& Title 2001) and down-drafting (Langhans et al. 2002), and (II) non-magnetic BP, which move upwards with the respective granules (same two papers).  We separate both types of BP using a pattern recognition code, (Bovelet \\& Wiehr 2007) especially adapted to detect the inter-granular BP. We then select magnetic inter-granular structures using their G-band-to-continuum ratio and additionally their Ca\\,{\\small II}\\,H excess brightness. There is some controversy about these criteria: Rezaei et al. (2007) find that the Ca\\,{\\small II} excess cannot be considered a general indicator for magnetic flux concentrations. Berger et al. (1998) argue that the G-band brightness is not a sufficient criterion for the magnetic nature of a BP. De\\,Wijn et al. (2005) suggest the Ca\\,{\\small II} brightness to be a more realistic criterion. Their concern about the G-band criterion refers to Berger \\& Title (2001), who, however, only exclude that criterion for those BP that appear embedded in granules, i.e., the above family\\,II.  We avoid these, restricting our selection to {\\it small} structures, thus excluding extended granular features and, hence, non-magnetic local G-band enhancements embedded. Among the small-scale inter-granular structures (IgS), we remove those that populate the non magnetic branch in the G-band-to-continuum scatterplot. The remaining IgS are than sampled by a significant Ca\\,{\\small II} excess.  The powerful pattern recognition algorithm, followed by a selective sampling of small structures, elimination of non-magnetic ones by their G-band-to-continuum ratio, and classification of magnetic ones by the Ca\\,{\\small II} excess, qualifies our study as an {\\it alternate way} to investigate small-scale solar kilo-Gauss magnetic flux-concentrations at the highest possible spatial resolution and independent of interpreting polarimetric measurements.   \\begin{figure} \\centering \\includegraphics[width=0.95\\linewidth,angle=0]{9717fig1.eps} \\caption{Magnetic map of the solar disk on April 12, 2007 from MDI showing an extremely low level of solar activity. The $149\\arcsec\\,\\times\\,117\\arcsec$ FOV investigated in this paper is marked.} \\label{FigMagnetSun}% \\end{figure} ",
        "conclusions": "The segmentation and selection of small-scale inter-granular structures with a significant Ca\\,{\\small II}\\,H excess allows lower limits to be determined for the number density and fractional area coverage of inter-granular kilo-Gauss magnetic flux-concentrations. A magnetic calibration of each feature yields a mean (unsigned) magnetic flux density in a quiet solar region close to existing results. Our method offers an alternate approach to polarimetry, although it inevitably misses those magnetic structures that do not exhibit a Ca\\,{\\small II}\\,H excess and `elusive magnetic structures' from fluctuations in thermodynamic properties (which S\\'anchez Almeida, 2000, estimate to carry at least half of the solar magnetic flux).  In case the total flux in our disk center FOV were representative of the entire solar surface ($221\\,\\times\\,$FOV), the Sun would be covered by $>\\!0.6\\!\\times\\!10^6$ MIgS with a total flux of $>\\!4.1\\,\\times\\,10^{23}$\\,Mx. Considering also magnetic structures which are not covered by our selection, the Sun at activity minimum may well be covered by a few millions of MIgS with a total flux of $10^{24}$\\,Mx or even more.  Our method offers a quantitative investigation of highly resolved (reconstructed) intensity images, avoiding the ambiguities of the Stokes calibration procedure and allowing closely neighboring opposite polarities to be disentangled up to the spatial resolution achieved in two-dimensional brightness images."
    },
    "1208/1208.4399_arXiv.txt": {
        "abstract": "We analyze blazar jet apparent speeds and accelerations from the RDV series of astrometric and geodetic VLBI experiments. From these experiments, we have produced and analyzed 2753 global VLBI images of 68 sources at 8~GHz with a median beam size of 0.9~milliarcseconds (mas), and a median of 43 epochs per source. From this sample, we analyze the motions of 225 jet components in 66 sources. The distribution of the fastest measured apparent speed in each source has a median of 8.3$c$ and a maximum of 44$c$. Sources in the 2FGL {\\em Fermi} LAT catalog display higher apparent speeds than those that have not been detected. On average, components farther from the core in a given source have significantly higher apparent speeds than components closer to the core; for example, for a typical source, components at $\\sim3$~mas from the core ($\\sim15$~pc projected at $z\\sim0.5$) have apparent speeds about 50\\% higher than those of components at $\\sim1$~mas from the core ($\\sim5$~pc projected at $z\\sim0.5$). We measure accelerations of components in orthogonal directions parallel and perpendicular to their average velocity vector. Parallel accelerations have significantly larger magnitudes than perpendicular accelerations, implying observed accelerations are predominantly due to changes in the Lorentz factor (bulk or pattern) rather than projection effects from jet bending. Positive parallel accelerations are significantly more common than negative ones, so the Lorentz factor (bulk or pattern) tends to increase on the scales observed here. Observed parallel accelerations correspond to modest source frame increases in the bulk or pattern Lorentz factor. ",
        "introduction": "\\label{intro} The bulk outflow of material at high Lorentz factors \\footnote{The Lorentz factor $\\Gamma=1/(1-\\beta^2)^{1/2}$, where $\\beta$ is the velocity expressed as a fraction of the speed of light.} in collimated relativistic jets is a well-established property of powerful blazars. Such high Lorentz factors can be directly observed through the high-speed apparent motions of jet components in VLBI imaging (e.g., Lister et al. 2009b, hereafter L09), and they are also required to explain blazar spectral energy distributions (e.g., Hartman et al. 2001), gamma-ray time variability (e.g., Dondi \\& Ghisellini 1995), and high radio-core brightness temperatures (e.g., Tingay et al. 2001). These relativistic jets must be accelerated over some length scale between about $10^{3}$ gravitational radii from the central black hole and the parsec scale where they are directly observed with VLBI (e.g., Sikora 2005; Vlahakis \\& K\\\"{o}nigl 2004). Although observations of high Lorentz factor flows are well-established, the theoretical mechanism by which these outflows are accelerated, and the length scale over which it operates, is not completely understood.  In the general framework of magnetic jet acceleration in blazars (e.g., Sikora 2005), the energy of the flow begins as magnetic energy, or Poynting flux, which is then converted into bulk kinetic energy during an acceleration phase, and finally into particle kinetic energy at shocks, which can then be radiated away. Magnetic acceleration has been investigated through general relativistic magnetohydrodynamic simulations (e.g., McKinney 2006); also see Komissarov (2011) and K\\\"{o}nigl (2010) for summaries of the theory of magnetic acceleration of relativistic jets.  A number of details within this general framework remain to be addressed: such as whether the jet is produced in a steady-state or whether the acceleration is impulsive (Granot et al. 2011; Lyutikov \\& Lister 2010), and whether the acceleration is complete before the scales observed with VLBI, or is still occurring on these parsec scales. Vlahakis \\& K\\\"{o}nigl (2004) argue that magnetic acceleration can continue to act out to the parsec scales, and they interpreted two specific observed acceleration events in NGC~6251 and 3C~345 as evidence for this ``magnetic driving'' on parsec scales, but at the time of that paper there were insufficient VLBI observations to address the question of whether acceleration on parsec scales was a common property of blazar jets in large samples. Even if parsec-scale acceleration events are observed to be commonplace, it does not necessarily prove the direct observation of conversion of Poynting flux to kinetic energy, since there may also be hydrodynamic means to produce accelerations in matter-dominated jets (e.g., Daly \\& Marscher 1988; Kadler et al. 2008), or the observations may be showing an increase in the Lorentz factors of patterns in the underlying flow.  Direct observations of intrinsic acceleration through VLBI imaging are difficult. Precise measurements of component positions at many epochs are needed to reliably measure a second derivative in a position versus time plot. For individual jet components, the apparent speed is given by the well-known formula: \\begin{equation} \\label{speedeqn} \\beta_\\mathrm{app}=\\frac{\\beta\\sin\\theta}{1-\\beta\\cos\\theta}, \\end{equation} where $\\beta c$ is the intrinsic speed and $\\theta$ is the angle of the motion to the line of sight. When observed in a single component then, changes in the apparent speed can be produced either from a change in the intrinsic speed or the viewing angle. Observations of many apparently accelerating components are needed to statistically distinguish between these two cases. In practice then, observations of many sources at many epochs, totaling thousands of VLBI images, are needed to measure variations in intrinsic speeds. While observations of either individual apparent component accelerations (e.g., Unwin et al. 1997; Homan et al. 2003) or apparent component accelerations in smaller samples of blazars (e.g., Homan et al. 2001; Jorstad et al. 2005) have been previously noted, the MOJAVE survey with 2424 total images was the first to investigate blazar jet accelerations through a large statistical sample (Homan et al. 2009, hereafter H09).  In this paper we present a continuation of our blazar jet kinematics study from Piner et al. (2007), hereafter Paper I, that is designed to enable measurements of blazar jet accelerations using the RDV (Research \\& Development -- VLBA) series of experiments on the VLBA (Petrov et al. 2009). The RDV series of experiments is observed primarily for the purposes of astrometry and geodesy, but because the experiments have occurred roughly every two months since the VLBA opened, and produce quality images, they are also useful for blazar astrophysics (e.g., Kovalev et al. 2008; Pushkarev \\& Kovalev 2012). This is only the second large-scale study of accelerations in the apparent motions of extragalactic jets (following H09).  In Paper I, we analyzed jet kinematics using 19 VLBI experiments observed over a 5 year time baseline from 1994 to 1998 (RDVs 1 through 10 and 12, plus 8 similar VLBI experiments that were conducted on the VLBA before the RDV series began). In that paper, we studied all sources that had been observed at 3 or more epochs over those 19 experiments, yielding a total of 966 images of 87 sources, which were used to measure apparent jet speeds.  In this paper, we expand on our study from Paper I by extending the analysis to a total of 50 VLBI experiments over a 10 year time baseline from 1994 to 2003 (adding the 31 new experiments RDVs 11 and 13 through 42), and studying the kinematics of all sources that have been observed at 20 or more epochs over those 50 experiments. This survey is hereafter referred to as the RDV survey: it now comprises 2753 VLBI images of 68 sources, with a median of 43 epochs of observation per source. The number of images is approximately tripled compared to Paper I, and slightly exceeds the 2424 images in the MOJAVE survey (Lister et al. 2009a). Note also that the maximum number of epochs per source from Paper I (19) is now less then the minimum number of epochs per source considered in this paper (20).  The RDV experiments have continued to the present; the most recent available at this writing is RDV~93 observed on 28 June 2012. Thus, there are already an additional 51 RDV experiments in the VLBA archive above what is included in this paper. If these additional experiments are completely imaged and model fit, they have the potential to approximately double the RDV survey size compared to what is included in this paper: to approximately 6000 total images and approximately 100 epochs per source. At the present time, imaging of the RDV experiments is continuing so that studies such as those presented in this paper could be extended in the future.  The organization of this paper is as follows: in $\\S$~\\ref{sample} we describe our sample selection. In $\\S$~\\ref{models} we describe the VLBI imaging and model fitting, and present a large table of Gaussian model components. In $\\S$~\\ref{speeds} we present our measurement of component speeds, and in $\\S$~\\ref{acc} our measurements of component accelerations. In $\\S$~\\ref{discussion} we discuss the physical implications of these results, and in $\\S$~\\ref{conclusions} we present our major conclusions. Throughout the paper, we assume cosmological parameters of $H_{0}=71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{m}=0.27$, and $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "\\label{conclusions} We studied the parsec-scale kinematics of a sample of 68 extragalactic jets using global VLBI observations at 8~GHz from the RDV experiment series, significantly expanding upon our previous such study from Paper I. We included in this study all sources observed at 20 or more epochs during a series of 50 VLBI experiments from 1994 to 2003. We produced and analyzed 2753 VLBI images from these experiments, with a median of 43 epochs of observation per source. In terms of angular resolution and temporal coverage, this RDV survey is similar to the MOJAVE survey (L09; H09). We fit Gaussian models to the visibilities associated with each image, and identified a total of 225 jet components in 66 sources that could be followed from epoch to epoch. Second-order polynomials were fit to $x(t)$ and $y(t)$ for each component to study its velocity and acceleration. Observational results related to the measured apparent speeds can be summarized as follows: \\begin{enumerate} \\item{When multiple moving components are present in a jet, components farther from the core tend (about 75\\% of the time) to have larger apparent speeds than components closer to the core, with high statistical significance.} \\item{The variation in apparent speeds from component to component within a source is significantly less than the variation in apparent speeds from source to source within the sample, showing the existence of a characteristic speed associated with each source.} \\item{The distribution of the fastest measured apparent speed in each source shows a maximum of 44$c$ and a median of 8.3$c$.} \\item{Sources detected by the {\\em Fermi} LAT gamma-ray telescope display higher apparent speeds, with a median of $12.4c$, than those that have not been detected, which have a median of $5.7c$.} \\item{Apparently stationary or slowly moving Low Pattern Speed (LPS) components are found in 19 sources. These LPS components are clustered within $\\sim4$~pc projected from the core, and may represent truly stationary features such as recollimation shocks.} \\end{enumerate}  We identified high-quality subsamples of the full set of 225 components for acceleration analysis, and for each of these components we analyzed the relative acceleration both parallel and perpendicular to the direction of the average velocity vector, as well as the difference between the direction of the average velocity vector and the weighted mean position angle. Observational results related to the measured accelerations and non-radial motions can be summarized as follows: \\begin{enumerate} \\item{Significant non-radial motion is common, occurring in about half of the components at $\\ge 2\\sigma$ significance. When non-radial motion occurs, it tends to align the component with the downstream jet structure.} \\item{`High' relative accelerations (magnitudes that are at least $2\\sigma$ above 0.1~yr$^{-1}$) are fairly common, and comprise about 1/4 of the parallel accelerations and 1/7 of the perpendicular accelerations, the same rates of high accelerations found by the MOJAVE survey (H09).} \\item{The distributions of relative parallel and perpendicular accelerations are statistically distinct, with the parallel accelerations having a larger average magnitude by a factor of about 1.7. This difference implies that there are intrinsic changes in component Lorentz factors that dominate over jet bending in producing the observed parallel accelerations.} \\item{Positive parallel accelerations statistically dominate over negative parallel accelerations. The weighted mean relative parallel acceleration for the 64-component subsample is 0.133$\\pm$0.014~yr$^{-1}$. A typical observed relative parallel acceleration of 0.1~yr$^{-1}$ corresponds to an increase in the bulk or pattern Lorentz factor in the reference frame of the host galaxy of order $\\dot{\\Gamma}/\\Gamma\\sim10^{-3}$~yr$^{-1}$ at a distance of order 100 parsecs (de-projected) from the core.} \\end{enumerate}  In summary, blazar jets each have a characteristic speed within about 100 parsecs (de-projected) of the supermassive black hole, with any sideways motions tending to move components down a channel in the direction of the previous component. An average component increases its apparent speed at an observed rate of about 10\\% per year at distances of about 100 parsecs (de-projected) from the core. This apparent acceleration corresponds to an increase in the Lorentz factor at a rate of about one part in 10$^{3}$ per year in the reference frame of the host galaxy. A minority of components have an apparent deceleration at these distances.  All of the above conclusions are statistical in nature, and will not necessarily apply to any particular individual source. When taken together with the similar kinematic results from the MOJAVE survey by H09, the acceleration results reported here show that modest changes in bulk or pattern Lorentz factors on parsec scales are a relatively common feature of relativistic jets. This observational result has now been confirmed in a mutually consistent manner with high statistical significance by two large VLBI surveys (although note that these two surveys are not completely statistically independent, because they have 37 sources in common). These observations are consistent with modest increases in the bulk kinetic energy on parsec-scales, although the source of this energy is not determined from these observations.  This paper and Paper I represent the tip of the iceberg of the astrophysics that can be done with the RDV data. With $\\sim100$ experiments observed to date, the total number of potential images is approximately 10,000 each at 8 and 2~GHz. For example, adding the $\\sim50$ experiments that have been observed since 2003 to this study could double the size of the kinematic survey presented here. Many things can also be studied other than kinematics; including flux variability and multiwavelength correlations, spectral index and core opacity (e.g., Kovalev et al. 2008; Pushkarev \\& Kovalev 2012), jet ridgelines and bending, and transverse structures in jets.  \\vspace{-0.1in}"
    },
    "1208/1208.1269_arXiv.txt": {
        "abstract": "{The properties of unresolved protostars and their local environment (e.g., disk, envelope and outflow characteristics) are frequently inferred from spectral energy distributions (SEDs) through comparison with idealized model SEDs. However, if it is not possible to image a source and its environment directly, it is difficult to constrain and evaluate the accuracy of these derived properties. In this proceeding, I present a brief overview of the reliability of SED modeling by analyzing dust continuum synthetic observations of realistic simulations.} ",
        "introduction": "\\label{sec:1}  Forming stars may be heavily obscured by their natal dust and gas, which inhibits direct imaging and causes source radiation to be reprocessed from shorter to longer wavelengths. The details of the multi-wavelength emission, i.e., the spectral energy distribution (SED), thus provide important indirect clues about the protostellar properties and environment. For example, the absence of $\\le10~\\mu$m emission usually signifies a very young source with a dense gas envelope; low or non-existent millimeter emission indicates a more evolved source, which has accreted or dispelled its envelope (e.g., [2]).  The information implicit in the reprocessed emission is commonly extracted by comparing the observed SED with idealized models of the protostellar source and gas distribution % that are post-processed with a radiative transfer code to produce SEDs. Input models that reproduce the observed SED provide good candidate representations of the underlying source properties. These can provide a wealth of physical details (e.g., source mass, disk mass and radius, envelope density and radius, outflow cavity size, inclination) that would otherwise be impossible to obtain with observations. However, a number of caveats complicate this technique, including degeneracy between parameters, adoption of symmetry, and assumption that the observed SED represents a single source rather than a multiple system or even a small cluster [3].  Using simulations, which have completely known source and gas information, it is possible to assess the accuracy of this method applied to unresolved observed sources. In this proceeding, we summarize the results of [6], who present a comparison between the true properties of sources within simulations of a turbulent, star formatting cloud and the properties inferred from synthetic SEDs. ",
        "conclusions": ""
    },
    "1208/1208.5356_arXiv.txt": {
        "abstract": "The Cherenkov Telescope Array (CTA) is a project for a next-generation observatory for very high energy (GeV--TeV) ground-based gamma-ray astronomy, currently in its design phase, and foreseen to be operative a few years from now. Several tens of telescopes of 2--3 different sizes, distributed over a large area, will allow for a sensitivity about a factor 10 better than current instruments such as H.E.S.S, MAGIC and VERITAS, an energy coverage from a few tens of GeV to several tens of TeV, and a field of view of up to 10~deg.  In the following study, we investigate the prospects for CTA to study several science questions that can profoundly influence our current knowledge of fundamental physics. Based on conservative assumptions for the performance of the different CTA telescope configurations currently under discussion, we employ a Monte Carlo based approach to evaluate the prospects for detection and characterisation of new physics with the array.  \\noindent First, we discuss CTA prospects for cold dark matter searches, following different observational strategies: in dwarf satellite galaxies of the Milky Way, which are virtually void of astrophysical background and have a relatively well known dark matter density; in the region close to the Galactic Centre, where the dark matter density is expected to be large while the astrophysical background due to the Galactic Centre can be excluded; and in clusters of galaxies, where the intrinsic flux may be boosted significantly by the large number of halo substructures. The possible search for spatial signatures, facilitated by the larger field of view of CTA, is also discussed. Next we consider searches for axion-like particles which, besides being possible candidates for dark matter may also explain the unexpectedly low absorption by extragalactic background light  of gamma-rays from very distant blazars. We establish the axion mass range CTA could probe through observation of long-lasting flares in distant sources. Simulated light-curves of flaring sources are also used to determine the sensitivity to violations of Lorentz Invariance by detection of the possible delay between the arrival times of photons at different energies. Finally, we mention searches for other exotic physics with CTA. ",
        "introduction": "} A major open question for modern physics is the nature of the dark matter (DM). There is a large body of evidence for the presence of an unknown form of gravitational mass, at scales from kiloparsecs to megaparsecs, that cannot be accounted for by SM particles. The observation by the WMAP satellite~\\cite{Komatsu:2010fb} of the acoustic oscillations imprinted in the cosmic microwave background quantifies the DM component as contributing about 25\\% of the total energy budget of the Universe. Being dominant with respect to the baryonic component, which accounts for only about 4\\% of the total energy density, DM shaped the formation of cosmic structures. By comparing the galaxy distributions in large redshift galaxy surveys~\\cite{Reid:2009xm}, and through $N-$body simulations of structure formation~\\cite{Springel:2008b,Anderson:2010df,Springel:2008by}, it is inferred that the particles constituting the cosmological DM had to be moving non-relativistically at decoupling from thermal equilibrium in the early universe (`freeze-out'), in order to reproduce the observed large-scale structure in the Universe and hence the term ``cold DM'' (CDM). This observational evidence has led to the establishment of a concordance cosmological model, dubbed $\\Lambda$CDM~\\cite{Press:1974, Sheth:1999su,  Springel:2005nw}, although this paradigm is troubled by some experimental controversies~\\cite{Klypin:1999uc,Kravtsov:2009gi,deNaray:2011hy,Walker:2011zu,BoylanKolchin:2011de,BoylanKolchin:2011dk}.   One of the most popular scenarios for CDM is that of weakly interacting massive particles (WIMPs), which includes a large class of non-baryonic candidates with mass typically between a few tens of GeV and few TeV and an annihilation cross-section set by weak interactions \\cite[see, e.g., Refs.][]{Bertone:2005a,Feng:2010gw}. Natural WIMP candidates are found in proposed extensions of the SM, e.g. in Super-Symmetry (SUSY)~\\citep{Jungman:1996a, Martin:1998a}, but also Little Higgs~\\citep{Schmaltz:2005ky}, Universal Extra Dimensions~\\citep{Servant:2003a}, and Technicolor models~\\citep{Nussinov:1985xr,Chivukula:1989qb}, among others. Their present velocities are set by the gravitational potential in the Galactic halo at about a thousandth of the speed of light. WIMPs which were in thermal equilibrium in the early Universe would have a relic abundance varying inversely as their velocity-weighted annihilation cross-section (for pure $s-$wave annihilation): $\\Omega_{\\rm CDM}h^2=3\\times10^{-27}\\rm{cm}^3\\rm{s}^{-1}/\\left(\\sigma_{\\rm ann}v \\right)$~\\cite{Jungman:1996a}. Hence for a weak-scale cross-section $\\left(\\sigma_{\\rm ann}v\\right) = 3\\times10^{-26}\\rm{cm}^3\\rm{s}^{-1}$, they naturally have the required relic density $\\Omega_{\\rm CDM}h^2=0.113\\pm0.004$, where $h=0.704\\pm0.014$ is the Hubble parameter in units of $100$ km s$^{-1}$ Mpc$^{-1}$~\\cite{Komatsu:2010fb}. The ability of WIMPs to naturally yield the DM density from readily computed thermal processes in the early Universe without much fine tuning is sometimes termed the ``WIMP miracle''.   In some SUSY theories, a symmetry called `$R$-parity' prevents a too rapid proton-decay, and as a side-effect, also guarantees the stability of the lightest SUSY particle (LSP), which is thus a prime candidate for a WIMP. WIMPs can annihilate to SM particles, and have hadron or leptons in the final products of annihilation. Thus from cosmic DM annihilations, one can expect emission of neutrinos, charged cosmic rays, multi-frequency electromagnetic radiation from charged products, and prompt gamma-rays~\\cite{Colafrancesco:2005ji}.  The detection of these final state particles can help to identify DM --- this is termed ``indirect DM detection''. Gamma-rays are not deflected by cosmic magnetic fields, and thus trace back to their origin. Therefore, observation of a gamma-ray signal from cosmic targets where DM is expected could prove conclusive about its nature . \\\\  In the context of gamma-ray astronomy, the differential flux of gamma-rays from within a solid angle $\\Delta\\Omega$ around a given astronomical target where DM is expected, can be written as: \\begin{equation} \\label{eqnp} \\frac{\\mathrm{d}\\Phi(\\Delta\\Omega,E_{\\gamma})}{\\mathrm{d}E_{\\gamma}}\\,= {\\rm B_F}\\cdot\\frac{1}{4\\pi}\\,\\underbrace{\\frac{\\left(\\sigma_{\\rm ann} v\\right)}{2\\,m^2_\\chi}\\,\\sum_i{\\rm BR}_i\\frac{\\mathrm{d}N^i_{\\gamma}}{\\mathrm{d}E_{\\gamma}}}_{Particle\\,Physics}\\,\\cdot\\,\\underbrace{\\widetilde{J}(\\Delta\\Omega)}_{Astrophysics}\\, , \\end{equation} where $\\left( \\sigma_{\\rm ann} v\\right)$ is the annihilation cross-section (times the relative velocity of the two WIMPs), $\\sum_i{\\rm BR}_i\\,\\mathrm{d}N^i_{\\gamma}/\\mathrm{d}E_{\\gamma} = \\mathrm{d}N_{\\gamma}/\\mathrm{d}E_{\\gamma}$ is the photon flux per annihilation summed over all the possible annihilation channels $i$ with branching ratios ${\\rm BR}_i$, and $m_\\chi$ is the mass of the DM particle. The `astrophysical factor' $\\widetilde{J}$ is the integral over the line of sight (los) of the squared DM density and over the integration solid angle $\\Delta\\Omega$: \\begin{equation} \\label{eqn:jbar} \\widetilde{J} = \\int_{\\Delta\\Omega}d\\Omega \\int_{\\rm los}\\mathrm{d}s\\,\\rho^2(s,\\Omega). \\end{equation} The remaining term ${\\rm B_F}$ in Eq.~(\\ref{eqnp}) is the so-called `boost factor' which is a measure of our ignorance of intrinsic flux contributions that are not accounted for directly in the formula.  There are various known mechanisms for boosting the intrinsic flux, among which we mention the inclusion of subhalos, and the existence of a `Sommerfeld enhancement' of the cross-section at low velocity regimes in models where the DM particles interact via a new long-range force. All numerical $N-$body simulations of galactic halos have shown the presence of subhalos populating the host halo~\\citep[see, e.g., Refs.][]{Springel:2008b, Diemand:2008a}. Such density enhancements, if not spatially resolved, can contribute substantially to the expected gamma-ray flux from a given object. This effect is strongly dependent on the target: in dwarf spheroidal galaxies (dSphs) for example the boost factor is only of ${\\cal O}(1)$~\\citep{Pieri:2009, Abramowski:2010zzt}, whereas in galaxy clusters the boost can be spectacular, by up to a factor of several hundreds~\\cite{SanchezConde:2011ap, Pinzke:2011ek, Gao:2011rf}.  On the other hand, the Sommerfeld enhancement effect can significantly boost the DM annihilation cross-section~\\citep{Sommerfeld:1932,Lattanzi:2009a}. This non-relativistic effect arises when two DM particles interact in a long-range attractive potential, and results in a boost in gamma-ray flux which increases with decreasing relative velocity down to a saturation point which depends on the DM and mediator particle mass. This effect can enhance the annihilation cross-section by a few orders of magnitude~\\citep{Pieri:2009, Abramowski:2010zzt}.   The current generation of IACTs is actively searching for WIMP annihilation signals. dSphs are promising targets for DM annihilation detection being among the most DM dominated objects known and free from astrophysical background. Constraints on WIMP annihilation signals from dSphs have been reported towards Sagittarius, Canis Major, Sculptor and Carina by H.E.S.S.~\\citep{Aharonian:2007km, Aharonian:2008dm, Abramowski:2010zzt}, towards Draco, Willman 1 and Segue 1 by MAGIC~\\citep{Albert:2007xg, Aliu:2008ny, Aleksic:2011jx}, towards Draco, Ursa Minor, Bo\\\"{o}tes 1, Willman 1 and Segue~1 by VERITAS~\\citep{Acciari:2010a, Aliu:2012ga}, and again towards Draco and Ursa Minor by Whipple~\\citep{Wood:2008a}.  Nevertheless, the present instruments do not have the required sensitivity to reach the ``thermal'' value of the annihilation cross-section $\\left(\\sigma_{\\rm ann} v\\right) = 3\\times10^{-26}\\rm{cm}^3\\rm{s}^{-1}$. A search for a WIMP annihilation signal from the halo at angular distances between 0.3$^{\\circ}$ and 1.0$^{\\circ}$ from the Galactic Centre has also recently been performed using 112\\,h of H.E.S.S.~data \\cite{Abramowski:2011}. For WIMP masses well above the H.E.S.S.~energy threshold of 100\\,GeV, this analysis provides the currently most constraining limits on $\\left(\\sigma_{\\rm ann} v\\right)$ at the level of a few$\\times 10^{-25}\\,\\mbox{cm}^3\\mbox{s}^{-1}$. H.E.S.S., MAGIC and VERITAS have also observed some galaxy clusters, reporting detection of individual galaxies in the cluster, but only upper limits on any CR and DM associated emission~\\cite{Aharonian:2008uq, Aharonian:2009, Aleksic:2009ir, Acciari:2009uq, Aleksic:2011cp, Abramowski:2012}.  Even though IACT limits are weaker than those obtained from the Fermi-LAT satellite measurements in the GeV mass range~\\citep{Abdo:2010b, Abazajian:2010sq, Hutsi:2010ai, Ackermann:2011}, they complement the latter in the TeV mass range.  Gamma-ray line signatures can also be expected in the annihilation or decay of DM particles in space, e.g. into $\\gamma\\gamma$ or $Z^0\\gamma$. Such a signal would be readily distinguishable from astrophysical gamma-ray sources which typically produce continuous spectra~\\cite{Bringmann:2011ye}. A measurement carried out by H.E.S.S.~\\cite{Spengler:2011lines} using over 100~h of Galactic Centre observations and over 1000~h of extragalactic observations complements recent results obtained by Fermi-LAT~\\cite{Abdo:2010nc}, and together cover about 3 orders of magnitude in energy, from 10 GeV to 10~TeV.\\\\  In this contribution, we focus on the prospects for DM searches with CTA, which are expected to improve on the current generation of IACTs on the following basis: \\begin{list}{--}{\\itemsep=0pt} \\item the energy range will be extended, from a few tens of GeV to several tens of TeV. At low energies, this will allow overlap with the Fermi-LAT instrument, and will provide sensitivity to WIMPs with low masses. For WIMPs with mass larger than about 100 GeV, CTA will have higher sensitivity as our studies indicate \\citep{Funk:2012}. \\item the improved sensitivity in the entire energy range, compared to current instruments, will obviously improve the probability of detection, or even \\emph{identification} of DM, through the observation of spectral features, \\item the increased FOV (about 10 deg versus $2-5$ deg) with a much more homogeneous sensitivity, as well as the improved angular resolution, will allow for much more efficient searches for extended sources like galaxy clusters (Section~\\ref{sec:gc}) and spatial anisotropies (Section~\\ref{sec:spatial}), \\item finally, the improved energy resolution will allow much better sensitivity to the possible spectral feature in the DM-generated photon spectrum. While astrophysical sources show typically power-law spectra with steepening at high energies, DM spectra are universal and generically exhibit a rapid cut-off at the DM mass. For specific models, ``smoking gun'' spectral features can appear \\citep{Bringmann:2011ye}. The observation of a few identical such spectra from different sources will allow both precision determination of the mass of the WIMP and its annihilation cross-section. \\end{list}     For the following studies, in order to have a detection, we require $a)$ the number of excess events over the background larger than 10 in the signal region, $b)$ the ratio between the number of excess events and the background events larger than 3\\%, and $c)$ the significance of the detection computed following Eq.~(17) of Li\\&Ma \\citep{Li:1983fv}, to be larger than $5$. If not explicitly mentioned, we used a number of background-control regions set to 5 ($\\alpha$ = 0.2 in the Li\\&Ma notation), which is a conservative choice, given the fact that the large FOV of CTA may allow for $\\alpha<0.2$. In case of non detection within a certain observation time, we calculate integral upper limits following the methods described in \\citet{Rolke:2005a} (bounded profile likelihood ratio statistic with Gaussian background, and with a confidence level of 95\\% C.L) in all cases expect the Galactic halo case, where we use the method of \\citet{FeldmanCousins1998}.  We study the effect of various annihilation spectra, assuming in turn 100\\% BR into a specific channel ($b\\bar{b}$, $\\tau^+\\tau^-$ or $\\mu^+\\mu^-$). The spectral shapes are obtained from different parameterisation from the literature \\citep{Tasitsiomi2002,Cirelli:2010xx,Cembranos2011}. For the $b\\bar{b}$ channel, which is used for comparison of different targets (see Fig.~\\ref{fig:Fermi.vs.CTA}), this difference accounts for few percents (depending on the DM mass), which is substantially smaller than the uncertainties in, e.g.,  the astrophysical factor, and do not significantly alters the conclusions.      \\subsection{Observations of dwarf satellite galaxies\\label{sec:dwarf}} In the $\\Lambda$CDM paradigm, galaxies such as ours are the result of a complex merger history and are expected to have extended halos of DM in accordance with observations. dSphs are satellites orbiting the Milky Way under its gravitational influence and are considered as privileged targets for DM searches for the following reasons: \\begin{list}{--}{\\itemsep=0pt} \\item the study of stellar dynamics shows that dSphs are among the most DM-dominated systems in the Universe, with mass-to-light ratio up to a few hundreds. In particular, the otherwise very uncertain astrophysical factor (Eq.~\\ref{eqn:jbar}) can be constrained by dynamical arguments \\citep{Evans:2003sc}, \\item many of the dSphs lie within $\\sim 100$ kpc  of the Earth, \\item they have favourable low gamma-ray backgrounds due to the lack of recent star formation history and little or no gas to serve as target material for cosmic-rays~\\citep{Mateo:1998wg}. \\end{list}   The family of dSphs is divided into ``classical'' dSphs, which are well-established sources with relatively high surface brightness and hundreds of member stars identified~\\cite{Simon:2007dq, Charbonnier:2011ft}, and ``ultra-faint'' dSphs, which have mainly been discovered recently through photometric observations in the Sloan Digital Sky Survey (SDSS)~\\cite{York:2000gk} and have very low surface brightness and only a few tens or hundreds of member stars. Some of the ultra-faint dSphs are not well-established as such because of similarity of their properties with globular clusters, hence their nature is often under debate. However, they are of particular interest due to their potentially very large, albeit uncertain, mass-to-light ratios.   \\begin{table}[!t] \\centering \\small{% \\begin{tabular}{l|ccc|l|c} \\hline dSph & Dec.  & D     & $\\widetilde{J}$         & Profile & Ref.\\\\ & [deg] & [kpc] & [GeV$^{2}$\\,cm$^{-5}$] &  \\\\ \\hline Ursa Minor & $+44.8$ & $66$ & $2.2\\times10^{18}$ & NFW & \\cite{Charbonnier:2011ft}\\\\ Draco      & $+34.7$ & $87$ & $7.1\\times10^{17}$ & NFW & \\cite{Charbonnier:2011ft}\\\\ Sculptor   & $-83.2$ & $79$ & $8.9\\times10^{17}$ & NFW & \\cite{Charbonnier:2011ft}\\\\ &         &      & $2.7\\times10^{17}$ & ISO & \\cite{Battaglia:2008jz}\\\\ Carina     & $-22,2$ & $101$& $2.8\\times10^{17}$ & NFW & \\cite{Charbonnier:2011ft}\\\\ \\hline Segue~1    & $+16.1$ & $23$ & $1.7\\times10^{19}$ & Einasto & \\cite{Aleksic:2011jx}\\\\ Willman~1  & $+51.1$ & $38$ & $8.4\\times10^{18}$  & NFW & \\cite{Acciari:2010a}\\\\ Coma Berenices & $+23.6$ & 44 & $3.9\\times10^{18}$ & NFW & \\cite{Strigari:2008} \\\\  \\hline \\end{tabular} } \\caption{\\label{tab:jbar} Astrophysical factors for a selection of the most promising classical and ultra-faint dSphs. Dec. is the target declination, D the distance and $\\widetilde{J}$ is defined as in Eq.~\\ref{eqn:jbar}.} \\end{table}  Table~\\ref{tab:jbar} shows the astrophysical factor $\\widetilde{J}$ for few selected dSphs for comparison. For the classical dSphs, we selected the two most promising Northern (Ursa Minor and Draco) and Southern (Sculptor and Carina) ones according to~\\citet[Table~2]{Charbonnier:2011ft}. The statistical uncertainties on the astrophisical factor are roughly one order of magntiude at 68\\% CL, slightly depending on the dSphs, and can be found in \\citep[Table~2]{Charbonnier:2011ft}. For the ultra-faint dSphs, we include Segue~1, Willman~1 and Coma Berenices, which have the highest $\\widetilde{J}$-values (although their nature is still under debate, especially for Segue~1~\\cite{Belokurov:2007a, NiedersteOstholt:2009a, Geha:2009a, XiangGruess:2009a, Simon:2010a, Essig:2009a, Martinez:2009jh}, which makes the determination of the astrophysical factor less accurate than for classical dSphs). We remark how the estimation of the astrophysical factor is subject to uncertainties of either statistical origin or due to the different assumptions considered for its calculation. A systematic study has been done for Sculptor, to estimate the effect of the profile shape and velocity anisotropy assumptions \\cite{Battaglia:2008jz}. Another compilation of astrophysical factors for several dSphs can be found in \\cite{Ackermann:2011}.  For the subsequent discussion, we consider only three sources: Ursa Minor and Sculptor representative of classic dSphs and located in the Northern and Southern hemisphere respectively, and Segue~1 having the largest astrophysical factor.  \\subsection*{Bounds on the annihilation cross-section} \\begin{figure}[h!t] \\includegraphics[width=0.95\\linewidth]{./Sculptor_NFWISO_Sub_100h_0.1_test.eps} \\caption{\\label{fig:profiles} CTA sensitivities on the velocity-averaged annihilation cross-section as a function of the WIMP mass for 100~hours observation of Sculptor with the CTA array $E$ (solid line), $B$ (dashed line) and $C$ (dashed-dotted line). Both the NFW (black line) and cored isothermal (ISO, red line) DM halo profiles are shown, for an integration solid angle $\\rm \\Delta\\Omega = 1\\times10^{-5}\\,sr$. Annihilations are assumed to occur with 100\\% branching ratio into $b\\bar{b}$.} \\end{figure}  Two kinds of radial profiles are generally used to model the DM distribution in dSphs: \\emph{cusped} and \\emph{cored} profiles~\\cite{Walker:2009zp}. While the former is motivated by numerical $N-$body simulations, the latter seems to be more consistent with observations~\\cite{Salucci:2011ee}, but the issue is still under debate ~\\citep[see, e.g.,][]{Valenzuela:2005dh}. The standard cusped profile is the Navarro, Frenk \\& White form (NFW)~\\cite{Navarro:1997a}, while more recently it has been shown that the Einasto profile~\\cite{Navarro:2010a} provides also a good fit to the subhalos in $N-$body simulations~\\citep{Springel:2008b}. On the other hand, for systems of the size of dSphs, the possibility of centrally cored profiles has also been suggested~\\citep{Moore:1994yx, Flores:1994gz,Walker:2011zu}. In conclusion, observations of low surface brightness and dSphs~\\citep{deBlok:2001fe,vandenBosch:2000rza,deBlok:2009sp} show that both cusped and cored profiles can accommodate their stellar dynamics.  Fig.~\\ref{fig:profiles} shows the integral upper limits towards Sculptor, the best Southern candidate from Table~\\ref{tab:jbar}, for which we consider both a cusped NFW~\\cite{Charbonnier:2011ft} and a cored isothermal~\\cite{Abramowski:2010zzt} profile. The sensitivity is calculated assuming that the DM particle annihilates purely in the $b\\bar{b}$ channel, for arrays $B,\\;C$ and $E$. The observation time is set to 100 hours and the integration solid angle to $\\Delta\\Omega = 10^{-5}$~sr.  The best reached sensitivity is at the order of few$\\times10^{-23}$~cm$^3$~s$^{-1}$ for the NFW profile for both arrays $E$ and $B$, while the isothermal profile is less constraining. Weaker constraints in the low mass range are obtained for the $C$ array due to the lack of the large-size telescopes in the centre of their layout. The capability of CTA to discriminate between the two profiles is therefore restricted.  The integration solid angle plays a central role in the estimation of the sensitivity and in the discrimination of the cusp or core profiles. The former point was addressed already \\cite[Fig.~7]{Charbonnier:2011ft} where it was shown that small integration angles guarantee the strongest constraints. In the case of CTA, depending on the array layout (and the energy range), the angular resolution could be as low as 0.02~deg, corresponding to a minimum integration angle of about $10^{-6}$~sr, and thus our results can be considered conservative, with an expected improvement of up to a factor $\\sim2$. Concerning the second point, \\citet{Walker:2011fs} showed that the more robust constraints, regardless of whether the profile is cored or cusped, are reached for an integration angle $r_c=2\\,r_{1/2}/D$, where $r_{1/2}$ is the so-called half-light radius, and $D$ is the distance to the dSph. For Sculptor, $r_c=0.52$, which is over 5 times the integration angle adopted here. In our calculation this would imply a weakening of the upper limits of a factor of a few.\\\\   \\begin{figure}[h!t] \\includegraphics[width=0.95\\linewidth]{./Sculptor_Segue1_UrsaMinor_BR_100h_0.1_test.eps} \\caption{\\label{fig:sigmav_ul} CTA sensitivities on the velocity-averaged annihilation cross-section versus the WIMP mass for 100~hours observation towards Sculptor, Ursa Minor and Segue~1, assuming 100\\% branching ratio into $b\\bar{b}$ (for Segue~1 also into $\\tau^+\\tau^-$ and $\\mu^+\\mu^-$). The calculations are done for array $E$ and $\\rm \\Delta\\Omega = 1\\times10^{-5}\\,sr$. } \\end{figure}  In Fig.~\\ref{fig:sigmav_ul} we show the integral upper limits for two classical dSphs, namely Ursa Minor and Sculptor in the Northern and Southern hemispheres respectively, as well as the ultra-faint dSph Segue~1. In order to span the variety of DM particle models, we study the effect of various annihilation spectra (computed using Ref.~\\cite{Cirelli:2010xx}), assuming in turn 100\\% BR into $b\\bar{b}$, $\\tau^+\\tau^-$ and $\\mu^+\\mu^-$ channels for the array $E$ and an observation time $\\rm T_{\\rm obs} = 100$~h. Assuming the annihilation to be purely into $\\tau^+\\tau^-$, the sensitivity reaches few $\\times10^{-25}$~cm$^3$~s$^{-1}$ for 100~h observation time of Segue~1. In comparing the different dSphs (assuming the reference annihilation channel $b\\bar{b}$) we see that even the most promising classical dSphs are less constraining than Segue~1 by over a factor of 10. However the uncertainties in the estimation of astrophysical factors for ultra-faint dSphs mean that this conclusion may not be reliable. Note that in the above calculations we did not assume any intrinsic flux boost factor, i.e. ${\\rm B_F}=1$ in Eq.~(\\ref{eqnp}).  \\subsection*{Bounds on Astrophysical factors and Boost factors} \\begin{figure}[h!t] \\includegraphics[width=0.95\\linewidth]{./J_mins_100h_new_color.eps} \\caption{ \\label{jfactors} The minimum value of the astrophysical factor required for a $5\\,\\sigma$ detection with $\\rm T_{\\rm obs} = 100\\,h$, versus WIMP mass. Two annihilation channels are considered for arrays $B,$ $C,$ and $E$: $b\\bar{b}$ (upper curves) and $\\tau^+\\tau^-$ (lower curves). The estimated astrophysical factor for Segue~1 is shown for comparison.} \\end{figure}  Another approach to estimate the capabilities of CTA for DM detection in dSphs consists in the evaluation of the statistical significance of the DM signal as a function of the DM particle mass $m_{\\chi}$ and the astrophysical factor, for different possible annihilation channels.  Hereafter, we calculate the minimum astrophysical factor $J_{\\rm min}$ required to reach a statistical significance of $5\\,\\sigma$ assuming an effective observation time of $100$~h, and the thermal cross-section $3 \\times 10^{-26} \\mbox{cm}^{3}\\mbox{s}^{-1}$.  This is shown in Fig.~\\ref{jfactors} for two annihilation channels: $b\\bar{b}$ (upper curves) and $~\\tau^+\\tau^-$ (lower curves), using analytical fits from Ref.~\\cite{Cembranos:2010xd}.  Again, three proposed CTA configurations are studied: B, C, and E. In order to put these values into context, we note that the largest astrophysical factor $\\widetilde{J}$ for known dSphs is that of Segue~1 at $1.7\\times 10^{19}$~GeV$^{2}$~cm$^{-5}$~\\cite{Essig:2010em}. From the figure we see that array $B$ is the most constraining over the whole energy range. It is clear that for a detection, the astrophysical factor of the dSph needs to exceed $10^{21}$~GeV$^{2}$~cm$^{-5}$, which is only 1--2 orders of magnitude smaller than that of the Galactic Centre (see Section~\\ref{sec:halo}). While we may expect a few such objects in the Milky Way halo~\\cite{Pieri:2008a}, they ought to have already been detected and identified by Fermi-LAT. Although this has not happened, one can envisage DM subhalos with no associated dSph (or one not bright enough optically to be detected), and therefore such gamma-ray emitters may be hidden among the unidentified Fermi sources~\\cite{Nieto:2011sx}.\\\\  Another way to evaluate the prospects of DM detection is by means of the intrinsic flux \\emph{boost factor} term ${\\rm B_F}$ in Eq.~(\\ref{eqnp}).  The minimum ${\\rm B_F}$ is computed as the ratio of the minimum astrophysical factor $J_{\\rm min}$ which provides a 5$\\,\\sigma$ detection in 100~h of observation time with CTA, to the observational astrophysical factor $\\widetilde{J}$ from the DM modeling of the dSphs. Again, the thermal cross-section $3 \\times 10^{-26} \\mbox{cm}^{3}\\mbox{s}^{-1}$ is assumed. Fig.~\\ref{fig:Boost} shows the minimum ${\\rm B_F}$ for a 1~TeV DM particle annihilating into $\\tau^{+}\\tau^{-}$. $J_{\\rm obs}$ is calculated for a NFW profile for all the cases except Segue~1, where an Einasto profile is considered. Considering that the boost factor from subhalos in dSph is only of ${\\cal O}(1)$, CTA observations of dSphs will be more sensitive to scenarios where Sommerfeld enhancement is at play, which may instead boost the signal up to ${\\cal O}(1000)$.   \\begin{figure}[h!t] \\includegraphics[width=0.95\\linewidth]{./final_boost_100h.eps} \\caption{\\label{fig:Boost} Minimum boost factor required for a $5\\,\\sigma$ detection in 100~h by array $B$, for the dSphs in Table~\\ref{tab:jbar} and a 1~TeV WIMP annihilating into $\\tau^{+}\\tau^{-}$. The density profiles are taken to be NFW, except for Segue~1 where an Einasto profile has been assumed. The smallest boost required is ${\\rm B_F}=25$ for Segue~1.} \\end{figure}  \\subsection{Observations of Galaxy Clusters\\label{sec:gc}}  Within the standard $\\Lambda$CDM scenario, galaxy clusters, with masses around $10^{14}-10^{15}$~M$_\\odot$, are the largest gravitationally bound objects and the most recent structures to form~\\cite{Voit:2004ah}. They are complex objects, relevant for both cosmological and astrophysical studies, and for what concerns DM searches \\cite{Colafrancesco:2005ji, SanchezConde:2011ap, Pinzke:2011ek, Blasi:2007pm,Jeltema:2008vu, Pinzke:2010st,Ackermann:2010rg, Cuesta:2011, Zimmer:2011vy, Huang:2011dq, Han:2012au}. DM, in fact, is supposed to be the dominant component of the cluster mass budget, accounting for up to 80\\% of its mass (the other components are the galaxies and the gas of the intra-cluster medium (ICM)). This is why clusters have been considered as targets for the indirect detection of DM, with the possibility of detecting the gamma-rays produced in the annihilation (or decay) of DM particles in the halo of the cluster.  $N-$body simulations of halo formation and evolution have also proven that, while the majority of early-formed, small structures merge together giving shape to more massive objects, some of the subhalos survive and are still present in the ``host'' halo of larger objects. Theoretical models foresee a huge number of these substructures at all scales down to $10^{-11}-10^{-3}$~M$_\\odot$ \\cite{Bringmann:2009vf}. These subhalos have the effect of contributing to the total gamma-ray emission from DM annihilations, and they may have important consequences for DM indirect detection. This is especially true for galaxy clusters, where the intrinsic flux ``boost'' from subhalos can be of order $100-1000$, in particular compared to the case of dSphs, explored previously, where the subhalos boost should contribute only marginally. Despite the fact that, due to their vicinity, dSphs are usually considered as the best sources for DM indirect detection, thanks to the subhalos boost, some authors claim that galaxy clusters have prospects of DM detection better or at least as good as those of dSphs \\cite{SanchezConde:2011ap,Pinzke:2011ek, Gao:2011rf}.  On the other hand, in galaxy clusters, emission in the gamma-ray range is not only expected by DM annihilation. Clusters may host an Active Galaxy Nucleus (AGN, that appear as point-like sources at very high energies) and radio galaxies.  The case of the Perseus galaxy cluster, which has been observed by MAGIC during several campaigns in the last years, is emblematic: MAGIC detected both the central AGN NGC-1275~\\cite{Aleksic:Perseus} and the off-centreed head-tail radio galaxy IC~310~\\cite{Aleksic:2010xk}. Moreover gamma-rays are expected to be produced also from the interaction of cosmic rays (CRs) with the ICM \\cite{Pinzke:2010st, Blasi:1999aj,Miniati:2003ep,Pfrommer:2007sz, Ensslin:2010zh}. The physics of the acceleration of CRs (electrons and protons) is not completely understood, but plausible mechanisms can be shock acceleration during structure formation, or galactic winds driven by supernovae. CRs can also be injected into the ICM from radio galaxy jets/lobes. At the energies of interest here (above 10 GeV), CRs emit gamma-rays from the processes associated with the decay of the neutral and charged pions produced in the interaction of the CRs with the ICM ambient protons \\cite{Colafrancesco:1998us,Colafrancesco:2010kx}. Most importantly, such a contribution is usually found to be larger than the one predicted from DM annihilation. It thus represents an unavoidable source of background for DM searches in galaxy clusters. To date, the deep exposure performed with the MAGIC stereoscopic system of the Perseus cluster~\\cite{Aleksic:2011cp} placed the most stringent constraints from VHE gamma-rays observations regarding the maximum CRs-to-thermal pressure to $\\langle X_{CR}\\rangle<1-2$\\%.  The purpose of this section is to estimate the CTA potential to detect gamma-rays from DM annihilation in the halo of galaxy clusters. First, the CR-induced emission only will be considered. This component represents, by itself, an extremely interesting scientific case, at the same time being a background complicating the prospects of DM detection.  Afterward, the ideal case of a cluster whose emission is dominated by DM annihilation only will be treated. Finally, the combination of the two components distributed co-spatially will be discussed.  It should be noted here that gamma-ray emission from both DM annihilation and CRs is spatially extended, even though not always co-spatial. In particular, \\citet{SanchezConde:2011ap} proved that, for the case of DM, the contribution of subhalos is particularly relevant away from the halo centre, so that annihilations can still produce a significant amount of photons up to a distance of $1-2$ degrees from the centre.  This represents a problem for current Cherenkov Telescopes since their FOV is limited to $3-5$ degrees. CTA will overcome this limitation, having a FOV of up to 10~deg (at least above 1 TeV) and an almost flat sensitivity up to several degrees from the centre. It is reasonable to expect, therefore, that CTA will allow a step-change in capability in this important area.  In this study, we selected two benchmark galaxy clusters: Perseus and Fornax. Perseus has been chosen because it is considered that with the highest CR-induced photon yield but a low DM content, and Fornax for the opposite reason: it is considered the most promising galaxy cluster for DM searches~\\cite{SanchezConde:2011ap, Pinzke:2011ek}. We recall that Perseus is located in the Northern hemisphere, while Fornax is in the Southern hemisphere.  To study the prospects for CTA we use two Monte Carlo simulations of the instrument response functions and of the background rates for \\emph{extended sources}, for the case of array $B$ and array $E$, which we recall, are representatives of well-performing arrays at low energies (array $B$) and in the full energy range (array $E$). The MC simulations were developed explicitly for the analysis of extended sources so that all the relevant observables are computed throughout the entire FOV.  \\subsection*{Gamma-ray emission from cosmic-rays} \\label{sec:cr}  Gamma-ray emission due to the injection of CRs into the ICM of a galaxy cluster is proportional both to the density of the ICM and the density of CRs. For the present work, we refer to the hadronic CR model of Pinzke and collaborators \\cite{Pinzke:2011ek, Pinzke:2010st}, based on detailed hydrodynamic, high-resolution simulations of the evolution of galaxy clusters, since in these works we found detailed morphological information, essential to compute the CTA response. The CR surface brightness rapidly decreases with the distance from the centre of the halo, so that, in most cases, the total emission is contained in $0.2-0.3\\,r_{200}$, where $r_{200}$ is the projected virial radius of the cluster, where the local density equals 200 times the critical density (see, e.g., Fig.~14 of \\citet{Pinzke:2011ek}, from which we derive the surface brightness of the clusters we analyze). $r_{200}=1.9$~Mpc ($1.4^\\circ$) for Perseus and $0.96$~Mpc ($2.8^\\circ$) for Fornax~\\citep{SanchezConde:2011ap}. The energy spectrum of the model, in the energies of interests here (above 10~GeV), is a power-law with a slope of $-2.25$.  Since the emission region is extended in the sky, we first divide the FOV into a grid of pixels each 0.2 degrees wide, and then we define the region of interest (ROI), constituted by all the pixels within an angle $\\theta_{max}$ from the centre of the camera. We consider 15 values of energy threshold $E_i$ in logarithmic steps from 50 GeV to 50 TeV. With the theoretical gamma-ray emission and the instrument response, we are able to compute the predicted number of background ($N_i^{\\mbox{\\tiny{OFF}}}$) and signal events ($N_i$) above each $E_i$, in each bin of the ROI separately, and then we integrate over the entire ROI. The model of \\citet{Pinzke:2011ek} predicts a rather large gamma-ray flux for Perseus ($\\Phi_{\\mbox{\\tiny{CR}}}(>100$~GeV$) = 2.04\\times 10^{-11} \\mbox{cm}^{-2}\\mbox{s}^{-1}$), the largest among the galaxy clusters, and a smaller one for Fornax ($\\Phi_{\\mbox{\\tiny{CR}}}(>100$~GeV$)=1.5\\times 10^{-13} \\mbox{cm}^{-2}\\mbox{s}^{-1}$). Above the different energy thresholds $E_i$, we determine how many hours CTA will need to detect the sources.  We perform the calculation for the two CTA array $B$ and $E$ and for different ROI. We repeat the procedure 10 times for each energy threshold and average the results, in order to quantify the statistical fluctuations occurring when the number of events (both $N_i$ and $N_i^{\\mbox{\\tiny{OFF}}}$) are generated. The results are shown in Fig.~\\ref{fig:Perseus_CR}.  \\begin{figure}[h!t] \\centering \\includegraphics[width=0.95\\linewidth]{./CRonly_7nov.eps} \\caption{\\label{fig:Perseus_CR} Integration time required to have a 5$\\sigma$ detection (see text for details) of gamma-rays from CR-induced gamma-rays only according to the model of~\\citet{Pinzke:2011ek} for Perseus (lower curves, $\\Phi_{\\mbox{\\tiny{CR}}}(>100$~GeV$) = 2.04\\times 10^{-11} \\mbox{cm}^{-2}\\mbox{s}^{-1}$) and Fornax (upper curves, $\\Phi_{\\mbox{\\tiny{CR}}}(>100$~GeV$)=1.5\\times 10^{-13} \\mbox{cm}^{-2}\\mbox{s}^{-1}$). The integration time is shown as a function of the energy threshold, and over different ROI, for the case of array $B$ (dashed lines) and array $E$ (solid lines). The shaded regions indicate the $1\\sigma$ standard deviation among 10 different simulations.} \\end{figure}  If one assumes the CR-induced gamma-ray model by \\citet{Pinzke:2011ek}, CTA will detect such radiation from Perseus already in about $100$~h, a fact which will constitute an extraordinary scientific result by itself\\footnote{We underline that the upper limits obtained by the MAGIC experiment on Perseus~\\citep[Fig.~3]{Aleksic:2011cp} already constrain by about 20\\% the model predictions (the same used here), impling that the maximum CR acceleration efficiency is lower than 50\\% or, alternatively, the presence of non-neligible CR transport phenomena.}. The discovery could indeed be potentially close, opening up a completely new observation window on the Universe. We underline that there is an absolute lower limit for gamma-rays in the hadronic scenario for clusters with an observed radio halo: a stationary distribution of CR electrons loses all its energy to synchrotron radiation for strong magnetic fields, as those in the radio halo, and therefore the ratio of gamma-ray to synchrotron flux becomes independent of the spatial distribution of the CRs and the thermal gas. % For the Perseus cluster this lower limit is roughly a factor $3-4$ from the gamma-ray flux predicted by the CR model~\\cite[see Fig 3 in Ref.][]{Aleksic:2011cp}, hence CTA would, in the worst case scenario, require about $1,000$ hours of observation to completely rule out the hadronic models. Such large observation times can in principle be achieved either by, e.g., multi-annual observational campaigns. On the other hand, a non detection with CTA in a few hundred hours would seriously constrain the model and thus pose interesting challenges on the galaxy cluster physics. The situation is more pessimistic for our model of Fornax, which is out of reach for CTA.  We see that the exact value of the integration time depends on the energy threshold chosen for the analysis. The reason for this is the tradeoff between the gamma-ray efficiency at different energies (the effective area), the source intrinsic spectrum and the chosen ROI. Roughly 90\\% of the CR-induced emission is expected within about $0.1\\,r_{200}$ for Perseus, which corresponds to roughly $0.2^\\circ$. We checked that integrating larger ROI, more background than signal is included in the analysis, thus deteriorating the significance of the detection. This suggests that in realistic cases, the best ROI should be optimized. Finally, we also see that the prospects of detection are similar for both considered arrays, $B$ and $E$.  \\subsection*{Gamma-ray emission from Dark Matter annihilation} The gamma-ray brightness due to DM annihilations from a particular viewing angle in the sky is proportional to the DM density squared integrated along the line of sight, as shown in Eq.~(\\ref{eqnp}). In the case of galaxy clusters, the contribution of the smooth DM halo is boosted by the presence of DM subhalos. Recent $N$-body simulations of Milky Way-like halos \\cite{Springel:2008b, Anderson:2010df, Springel:2008by} found that the contribution of subhalos is small in the centre of the halo, due to dynamical friction and tidal effects that disrupt the subhalos. However, already at distances of $0.01-0.05\\,r_{200}$, subhalos become the dominant component. The real value of the boost factor from subhalos is unknown and the theoretical estimates depend on different assumptions and different methods used in the calculations. \\citet{Pinzke:2011ek} estimated a ${\\rm B_F}=580\\mbox{ and }910$ for Fornax and Perseus respectively (for a minimal halo mass of $10^{-6}$~M$_\\odot$) , while other authors gave ${\\rm B_F}$ from few tens \\cite{SanchezConde:2011ap} up to several thousands \\cite{Gao:2011rf}.  We refer again to the results of \\citet{Pinzke:2011ek} where the authors assumed a double power-law to describe the luminosity of subhalos as a function of the projected distance from the centre of the halo, a behavior derived by analyzing the sub-halos in the Aquarius $N$-body simulation. They also found the projected surface brightness to be largely independent of the initial profile of the smooth DM halo. As a result, the DM profile is very flat since the emission decreases approximately only $10$\\% at a distance of $1.5-2.0$ degrees from the centre, depending on the cluster~\\cite[Fig.~\\ref{fig:morpho} and Ref.][]{SanchezConde:2011ap}. For the case of Perseus and Fornax, we used the results of Fig.~10 of Ref.~\\cite{Pinzke:2011ek}, assuming a telescope angular resolution of $0.1$ degree, which is a good approximation for CTA, despite the fact that the exact value depends on the array, the energy and the position in the FOV. We underline that in the case of galaxy cluster, the contribution from substructure strongly shapes the region of emission, basically moving from a point-like source (in case no substructure are considered), to an extended source. Given the fact that the analysis differ in the two cases, the contribution from substructure cannot be considered as a simple multiplicative factor in the intrinsic expected flux with respect to point-like case.  Hereafter we consider the Fornax cluster, which has the largest expected DM-induced photon yield. The intrinsic flux is taken from \\cite[Table~2]{Pinzke:2011ek} and includes an intrinsic boost factor from subhalos of ${\\rm B_F}=580$, summing up to a total flux of $\\Phi_{\\mbox{\\tiny{DM}}}(>100$~GeV$)=3.6\\times 10^{-13} \\mbox{cm}^{-2}\\mbox{s}^{-1}$. Additional intrinsic boost factor may come from either other contributions from subhalos not accounted in this model, by mechanisms like the Sommerfeld enhancement discussed above, or by the effect of contraction processes due to baryonic condensation~\\cite{Gnedin:2004cx,Ando:2012vu}.  To compute the CTA prospects of detection, we consider only the case of DM annihilating into $b\\bar{b}$ (spectral shape obtained from Ref.~\\citep{Cembranos:2010xd}), while other channels like $\\tau^+\\tau^-$ or $\\mu^+\\mu^-$ may be more constraining, depending on the energy (see Fig.~\\ref{fig:sigmav_ul}). We take the reference thermal cross-section $3 \\times 10^{-26} \\mbox{cm}^{3}\\mbox{s}^{-1}$ and we scan DM particle mass $m_\\chi$ between 50~GeV and 4 TeV. We optimized the upper limit calculation as described in Ref.~\\cite{Aleksic:2011jx}, by optimizing the energy threshold above which the upper limit is estimated. In addition, we consider the possibility of extending the size of the ROI up to a $\\theta_{max}$ of 2 degrees, to encompass the full radial extension of the source. Fig.~\\ref{fig:Fornax_DM} show the results. In 100~h observation, the lack of detection would place exclusion limits at the level of $10^{-25}\\rm{cm}^{3}\\rm{s}^{-1}$.  We also studied the effect of integrating over larger and larger regions: despite the increased numbers of background events, the signal yield is also larger and, in the case of Fornax, we gain more in integrating up to $\\theta_{max}=1^\\circ$ than $0.5^\\circ$, while integrating over larger regions leads to a worse sensitivity.   \\begin{figure}[h!t] \\includegraphics[width=0.95\\linewidth]{./Fornax_final.eps} \\caption{\\label{fig:Fornax_DM} Prospect of detection of DM-induced signal from Fornax for a DM particle annihilating into $b\\bar{b}$ and 100~h integration time. The reference model is taken from Ref.~\\cite{Pinzke:2011ek} with subhalo boost factor ${\\rm B_F}=580$. The   shaded regions indicate the $1\\sigma$ standard deviation among 10 different simulations.} \\end{figure}   \\subsection*{Distinguishing the dark matter signal from other gamma-ray contributions}  \\begin{figure}[h!t] \\includegraphics[width=0.95\\linewidth]{./morpho.eps} \\caption{\\label{fig:morpho} The surface brightness (above 1 GeV) of the gamma-ray emission from the Fornax cluste from CRs (red), DM (blue) and the sum of the two contributions (black). The DM emission is calculated from the K' benchmark model of \\citep{Bringmann:2008kj} which has mass of 570~GeV and a velocity-averaged cross-section of $4.4\\times10^{-26}$ cm$^3$ s$^{-1}$. Adapted from Ref.~\\citep{Pinzke:2011ek}. } \\end{figure}  In the previous sections we have considered separately the contributions of CR and DM to the total gamma-ray photon yield. This is an unrealistic situation: galaxy clusters are, in fact, complex objects where gamma-rays may be due to different contributions possibly of different spatial origin: by collisions of accelerated CRs, by DM annihilations \\emph{and} by foreground or embedded astrophysical sources.  Fortunately, gamma-rays of different origin typically have different spectral shapes, with the DM-induced emission characterized by the peculiar cut-off at $E=m_\\chi$ and other remarkable spectral features \\cite{Bringmann:2007nk, Bringmann:2008kj}, in contrast to the plain spectral shapes (typically power-laws within the energy range of interest here) of the emission due to CRs, of the central galaxy or any astrophysical objects in the cluster. In the case a VHE emission is detected from a cluster, this fact may be used as a probe to discriminate between the components. However, we remark that in order to significantly discriminate the two sources one would need a quite significant detection over the CR-signal, which is often not supported by theoretical predictions for most galaxy clusters.  A distinct approach could be based on the different spatial extensions of the various contributions of VHE gamma-ray photons from galaxy clusters. The possible individual galaxies emitting within the cluster are typically seen as point-like sources, and thus one may exclude them from the FOV for CR and DM searches.  Moreover, from the fact that CR-induced radiation is more concentrated than that induced by DM, one can optimize the ROI to select only those where the emission is DM dominated. In Fig.~\\ref{fig:morpho}, we show the expected brightness profile for CR and DM photons for the Fornax cluster. One can see that up to $\\theta=0.4^\\circ$ the emission is dominated by CR-induced photons, whereas this exact value is cluster-dependent and model-dependent, and in particular the possible intrinsic boost-factor in the DM signal can affect this. In this example, above $\\theta=0.4^\\circ$, the CR-signal fades more rapidly than the DM one. Then, in principle, by considering a ROI with a $\\theta_{min}=0.4^\\circ$, one could be able to isolate the DM signal. The maximum integration angle $\\theta_{max}$ should be optimized according to the specific cluster and emission profile to maximize the sensitivity, as discussed above. Unfortunately, at the moment of writing this report, we did not have sufficient coverage in the MC of extended sources to perform such a study, and we are limited to a qualitative discussion. We mention that the ``geometrical'' discrimination makes sense only if the DM signal is sufficiently large, otherwise different observational strategies could be more constraining. Finally, we stress again that with a large FOV (at least above 1 TeV) that has a near constant sensitivity over several degrees will allow CTA to study extended high energy gamma-ray sources in detail for the first time, with possibly revolutionary consequences for the IACT technique.    \\subsection{Observations of the Galactic Halo and Centre\\label{sec:halo}}  The Galactic Centre (GC) is a long-discussed target for indirect DM searches with Cherenkov telescopes~\\citep{Gondolo:1999ef}. The density of the DM halo should be highest in the very centre of the Milky Way, giving rise to a gamma-ray flux from annihilation of DM particles. On the one hand, this view is strengthened by the results of recent $N-$body simulations of CDM halos \\citep{Aquarius2008} suggesting that, for an observer within the Milky Way, the annihilation signal from DM is not primarily due to small subhalos, but is dominated by the radiation produced by diffuse DM in the main halo. On the other hand, searches close to the GC are made difficult by the presence of the Galactic Centre source HESS J1745$-$290 \\citep{HessGC2010, HessGC2009} and of diffuse emission from the Galactic plane \\citep{HessRidge2006}. Both emissions can be plausibly explained by astrophysical emission processes: HESS J1745$-$290 is thought to be related to the Black Hole Sgr A$^{\\ast}$ or the pulsar wind nebula G 359.95$-$0.04 \\citep{GCPWN2006}, and the diffuse emission is well described as arising from hadronic cosmic rays interacting in giant molecular clouds. In both cases, the measured energy spectra do not fit DM model spectra \\citep{HessGC2006} and thus make a dominant contribution from DM annihilation or decay unlikely.  In this situation, DM searches should better target regions which are outside the Galactic plane and hence not polluted by astrophysical gamma-ray emission, but which are still close enough to the GC to exhibit a sizable gamma-ray flux from DM annihilation in the Milky Way halo~\\citep{Schwanke2009}. Given the angular resolution of Cherenkov telescopes and the scale height of the diffuse emission from the  Galactic plane these criteria are fulfilled for an angular distance of about $0.3^{\\circ}$ from the GC. This angular scale translates into a distance of 45\\,pc from GC when using 8.5\\,kpc as the galactocentric distance. The radial DM density profiles obtained in $N-$body simulations of Milky Way sized galaxies, like Aquarius \\citep{Aquarius2008} and Via Lactea II \\citep{ViaLactea2008}, can be described by Einasto and NFW parameterizations, respectively. These parameterizations differ substantially when extrapolating to the very centre of the Milky Way halo since the NFW profile is much more strongly peaked. At distances greater than about 10\\,pc, the difference is, however, just a factor of 2 which implies that a search at angular scales of $>0.3^{\\circ}$ will not be hampered by the imprecise knowledge of the DM density profile at small scales.  A search for a DM annihilation signal from the halo at angular distances between 0.3$^{\\circ}$ and 1.0$^{\\circ}$ from the GC has recently been performed using 112\\,h of H.E.S.S.~data \\citep{Abramowski:2011}. For WIMP masses well above the H.E.S.S.~energy threshold of 100\\,GeV this analysis provides the currently most constraining limits on the velocity averaged annihilation cross section $\\left(\\sigma_{\\rm ann} v\\right)$ of WIMPs (for IACTs) at the level of few $10^{-25}\\,\\mbox{cm}^3\\mbox{s}^{-1}$. Towards lower WIMP masses, observations of dwarf galaxies with the Fermi-LAT satellite yield even better limits \\citep{Abdo:2010b} demonstrating how both observations of dwarf galaxies and of the extended GC region allow to jointly constrain the parameter space.   \\subsection*{Simulations and Assumptions} The prospects of a search for DM annihilation photons from the Milky Way halo with CTA depend on $(i)$ the performance of the southern CTA array, $(ii)$ the applied analysis and background rejection techniques, and $(iii)$ the details of the DM distribution and WIMP annihilation. At low energies, the sensitivity of IACTs is limited by the presence of hadron and electron showers which arrive isotropically and which can only be distinguished from photons on a statistical basis. The basic strategy for the halo analysis is therefore to compare the fluxes of gamma-like events from a signal region (with solid angle $\\Delta\\Omega_s$) and a  background region (solid angle $\\Delta\\Omega_b$) and to search for DM features in the background-subtracted energy spectra. The signal region can be chosen such  that it has the same instrumental acceptance as the background region, but is located closer to the GC and features therefore a higher DM annihilation flux. For the purpose of this section, we rewrite Eq.~(\\ref{eqnp}) in terms of differential DM photon \\emph{rate} expected from the signal or background regions ($s,\\,b$ respectively), given by: \\begin{equation} \\frac{dR}{dE}|_{\\,\\mathrm{s,b}} = \\frac{\\left(\\sigma_{\\rm ann} v\\right)}{8\\pi\\,m^2_\\chi} \\frac{dN_{\\gamma}}{dE_{\\gamma}} \\int_{\\Delta\\Omega_{\\,\\mathrm{s,b}}} \\negspcfive J(\\Omega) A(\\Omega,E) d\\Omega\\mbox{,} \\end{equation} where $dN_{\\gamma}/dE_{\\gamma}$ is the photon spectrum generated in the annihilation of a WIMP of mass $m_\\chi$, and $A(\\Omega,E)$ are the CTA effective areas for photons, which depend on the position of the region within the FOV ($\\Omega$), the energy $E$ and further parameters (like the zenith angle of the observations). $J(\\Omega)$ is the line-of-sight integral over the squared DM density  $\\rho(r)$ (cf.~Eq.~\\ref{eqn:jbar}). Since the DM density depends only on the distance to the GC $r$ the line-of-sight integral and the astrophysical factor are only a function of the angular distance $\\psi$ from the GC. Assuming that the signal and background region differ only with respect to their DM annihilation flux and their relative size $\\alpha = \\Delta\\Omega_{\\mathrm{s}}/\\Delta\\Omega_{\\mathrm{b}}$, the rate of excess photon events  $R_{\\mathrm{s}} - \\alpha R_{\\mathrm{b}}$ is given by \\begin{equation} \\frac{\\left(\\sigma_{\\rm ann} v\\right)}{8\\pi m^2_\\chi} \\int_0^{\\infty}\\negspcfive dE\\frac{dN_{\\gamma}}{dE_{\\gamma}} \\left[ \\int_{\\Delta\\Omega_{\\mathrm{s}}}\\negspcfive J(\\psi) A(\\Omega,E) d\\Omega - \\alpha\\int_{\\Delta\\Omega_{\\mathrm{b}}}\\negspcfive J(\\psi) A(\\Omega,E) d\\Omega \\right] \\mbox{.} \\label{eq:halo_rate} \\end{equation}  Clearly, the rate vanishes when the astrophysical factors of the signal and the background regions are identical which implies that in the case of an isothermal DM density profile, a halo analysis with signal and background region chosen too close to the GC will not allow the placement of limits on $\\left(\\sigma_{\\rm ann} v\\right)$.  \\begin{figure}[h!t] \\centering \\includegraphics[width=0.75\\linewidth]{rings.eps} \\caption{\\label{fig:halo_rings}Illustration of the {\\em Ring Method} for constructing signal and background regions within one FOV of the CTA candidate arrays. The red star denotes the position of the GC in galactic coordinates; the blue star marks the pointing position of the CTA array which is shifted by an amount $b$ in latitude from the GC. The annulus with inner and outer radii $r_1$ and $r_2$ around the observation position defines regions of equal acceptance. The signal region (blue, close to the GC) is constructed as intersection of the annulus and a circle around the GC with radius  $\\Delta_{\\rm{cut}}$. The remaining regions on the annulus (red) are used as background region. Regions within $\\pm 0.3^{\\circ}$ of the galactic  plane (yellow) are neither part of the signal nor of the background region. } \\end{figure}  \\begin{table}[h!t] \\centering \\begin{tabular}{c|cccc} \\hline Array & $b$ & $r_1$ & $r_2$ & $\\Delta_{\\mathrm{cut}}$  \\\\ \\hline E & $1.42^{\\circ}$ & $0.55^{\\circ}$ & $2.88^{\\circ}$ & $1.36^{\\circ}$\\\\ B & $1.40^{\\circ}$ & $0.44^{\\circ}$ & $2.50^{\\circ}$ & $1.29^{\\circ}$\\\\ \\hline \\end{tabular} \\caption{\\label{tab:halo_ring_param}Optimized values of the parameters used in the application of the {\\em Ring Method} for the candidate arrays $E$ and $B$. See Fig.~\\ref{fig:halo_rings} for a description of the parameters.} \\end{table}  Given an observation time $T$, Eq.~\\ref{eq:halo_rate} can be used to estimate the number of excess photons for a particular realization of CTA and a DM model defining  $\\left(\\sigma_{\\rm ann} v\\right)$, $dN_{\\gamma}/dE_{\\gamma}$ and $J(\\psi)$. Equivalently, one can place a limit on $\\left(\\sigma_{\\rm ann} v\\right)$ given an upper limit on the number of excess photon events. Simulations of the candidate arrays $E$ and $B$ at a zenith angle of $20^{\\circ}$ were used to infer the effective area for diffuse photons and the residual rate of protons anywhere in the FOV. Both arrays feature large-size telescopes and are therefore suitable for studies in the low-energy domain.  The available observation time was set to 100\\,h, which is about 10\\,\\% of the total observation time per year.%  Two different ways of defining signal and background regions were employed and compared, namely the so-called {\\em Ring Method} and the On-Off Method. For the {\\em Ring Method}, the candidate arrays $E$ or $B$ were assumed to observe the GC region at Galactic longitude $l=0$ and Galactic altitude $b$, and signal and background regions were placed in the same FOV as illustrated in Fig.~\\ref{fig:halo_rings}. An annulus with inner radius $r_1$ and outer radius $r_2$ around the observation position was constructed and divided into signal and background region such that the signal region is closer to the GC and has therefore a larger astrophysical factor. The separation of signal and background region is achieved by a circle with radius $\\Delta_{\\mathrm{cut}}$ around the GC whose intersection with the annulus defines the signal region. All other regions on the ring were considered as background region. The values of the four parameters $b$, $r_1$, $r_2$ and  $\\Delta_{\\mathrm{cut}}$ were optimized such that the attained significance of a DM signal per square root time was maximized. The maximization was carried out for a wide range of WIMP masses but the dependence on the actual WIMP mass was found to be fairly weak. The derived values for both candidate arrays are listed in Tab.~\\ref{tab:halo_ring_param}. Judging from present IACT observations, we do not expect strong diffuse gamma-ray emission to extend outside the $\\pm0.3^\\circ$ box used to mask the galactic disc. New point-like or slightly extended sources will be excluded, making the On and Off region smaller. In addition, the approach is only sensitive to gradients in the diffuse gamma-ray emission, whereas the charged particle background is isotropic. In the optimization process an Einasto profile was assumed for the DM signal, but the optimal values are only weakly dependent on the assumed profile in the region beyond $0.3$ degrees from the Galactic plane.  The usage of the annulus implies the same acceptance for signal and background region since the acceptance is, to good approximation, only a function of the distance to the  observation position. Placing both signal and background regions in the same FOV implies that both regions will be affected by time-dependent effects in a similar way. A disadvantage is, however, that the angular distance between the signal and background region is only of order of the FOV diameter, reducing the contrast in Eq.~\\ref{eq:halo_rate} considerably. This contrast was increased in the On-Off Method where data-taking with an offset of typically 30' in Right Ascension was assumed. In this mode, the telescopes first track for half an hour the same observation position as in the {\\em Ring Method} which defines the signal region. The telescopes then slew back and follow the same path on the sky for another 30\\,min. The second pointing has the same acceptance as the first one since the same azimuth and zenith angles are covered but generates a background region with much increased angular distance to the GC. In the On-Off Method, the observation time was 50\\,h for the signal and 50\\,h for the background region giving again a total observation time of 100\\,h. Regardless of whether the {\\em Ring Method} or the On-Off Method was used, all areas with $|b|<0.3^{\\circ}$ were excluded from signal and background regions to avoid pollution from astrophysical  gamma-rays. \\begin{table}[h!t] \\centering \\begin{tabular}{l|cccc} \\hline &   &  & &  \\\\[-3.5mm] Method & Array & $m_\\chi$ &  $\\widetilde{J}_{\\mathrm{s}}$ & $\\Delta\\Omega$ \\\\ &       & (TeV)      &  ($10^{22}\\,\\mbox{GeV}^2\\,\\mbox{cm}^{-5}$) & (sr) \\\\ \\hline Ring    & E & any & 4.68 & 0.00117 \\\\ & B & any & 4.43 & 0.00104 \\\\ \\hline On-Off  & E & 0.1 & 16.4 & 0.00751 \\\\ &   & 1   & 19.7 & 0.01044 \\\\ &   & 10   & 28.7& 0.02211 \\\\ \\hline On-Off  & B & 0.1 & 16.4 & 0.00751 \\\\ &   & 1   & 22.8 & 0.01384 \\\\ &   & 10   & 28.7 &0.02211 \\\\ \\hline \\end{tabular} \\caption{\\label{tab:halo_jfactor}Astrophysical factor for the signal region and size of the integration region ($\\Delta\\Omega$) for {\\em Ring} and On-Off method. In the case of the On-Off Method, $\\Delta\\Omega$ was chosen as the entire FOV of the candidate array which introduces a dependence on the assumed WIMP mass since the effective FOV  grows with photon energy. The table gives values for a WIMP mass of 0.1, 1, and 10\\,TeV. } \\end{table}  The astrophysical factor (Eq.~\\ref{eqn:jbar}) was taken from the Aquarius Simulation~\\citep{Aquarius2008} which had been corrected for the presence of subhalos below the resolution limit of the simulation. The line-of-sight integral assumes a value of $40.3\\times 10^{24}\\,\\mbox{GeV}^2\\,\\mbox{cm}^{-5}\\,\\mbox{sr}^{-1}$ at $\\Psi=1^{\\circ}$. Table~\\ref{tab:halo_jfactor} lists the astrophysical factors of the signal regions which were defined in the {\\em Ring} and On-Off Method, respectively. In case of the On-Off Method, the signal region was defined as the total effective FOV of the On--pointing which introduces a dependence on the WIMP mass since the FOV grows with photon energy. For the WIMP annihilation spectrum $dN_{\\gamma}/dE_{\\gamma}$ several different choices were considered. The generic Tasitsiomi spectrum \\citep{Tasitsiomi2002} is appropriate for a dominant annihilation into quark-antiquark pairs with subsequent hadronization into $\\pi^0$ particles and was used in the optimization of the parameters of the {\\em Ring Method}. Other spectra were explored by considering $b\\bar{b}$, $\\tau^+\\tau^-$ and $\\mu^+\\mu^-$ final states \\citep{Cembranos:2010xd}.  \\subsection*{Discussion} \\begin{figure}[h!t] \\centering \\includegraphics[width=0.95\\linewidth]{methods.eps} \\includegraphics[width=0.95\\linewidth]{spectra.eps} \\caption{\\label{fig:halo_methods_spectra} CTA sensitivities on the velocity averaged annihilation cross-section as a function of the WIMP mass. Shown are curves for the candidate arrays $E$ (blue) and $B$ (red).  {\\bf Top:} Comparison of the {\\em Ring Method} (solid lines) and On-Off Method for background subtraction. Annihilation as in \\citet{Tasitsiomi2002} was assumed. {\\bf Bottom:} Comparison of different WIMP spectra for the {\\em Ring Method}. The solid line denotes the case of annihilation into $b\\bar{b}$; $\\mu^+\\mu^-$ and $\\tau^+\\tau^-$ spectra are shown by the dotted and dashed lines, respectively. On both panels, the classical annihilation cross section for thermally produced WIMPs at $3\\times10^{-26}\\,\\mbox{cm}^3\\mbox{s}^{-1}$ is indicated by the black horizontal line. } \\end{figure}  The two plots in Fig.~\\ref{fig:halo_methods_spectra} show the upper limits for WIMP masses between 0.1\\,TeV and 10\\,TeV, translated from the sensitivity using here the method of~\\citet{FeldmanCousins1998}.  Each curve corresponds to one set of assumptions. It is evident that the most constraining limits can be derived for masses of about 0.5\\,TeV which is a factor of 2 improvement compared to current IACT arrays like H.E.S.S. reaching best sensitivity around 1\\,TeV. This is a direct consequence of the lower threshold and superior stereoscopic background rejection of the CTA candidate arrays. Typical limits are around few $10^{-26}\\,\\mbox{cm}^3\\mbox{s}^{-1}$ which is a factor of 10 improvement compared to current IACTs. The comparison of array $E$ (blue) and $B$ (same line style but red) shows that the limits for array $B$ are always better, which can be understood from the fact that $B$ comprises $5$ large-size telescopes and array $E$ only $4$. The magnitude of this effect is, however, comparatively small ($\\sim 20\\,\\%$). Overall, CTA should be able to probe the parameter space below the velocity averaged annihilation cross-section for thermally produced DM of $3\\times 10^{-26}\\,\\mbox{cm}^3\\mbox{s}^{-1}$ for WIMP masses between several ten GeV and several TeV.  The upper panel of Fig.~\\ref{fig:halo_methods_spectra} illustrates the impact of data-taking with the {\\em Ring Method} and the On-Off Method for the case of a dominant annihilation into quark-antiquark pairs with subsequent $\\pi^0$ creation \\citep{Tasitsiomi2002}. The On-Off Method (dashed lines) is more sensitive than the {\\em Ring Method} (dashed lines). One must keep in mind, however, that the On-Off Method spends 50\\,\\% of the observation time far away from the GC which implies that this data set will be of limited use for studies of astrophysical sources. Another drawback of the On-Off Method is its susceptibility to systematic effects arising from variations of the data-taking conditions (electronics, atmosphere). In view of this, the increased sensitivity for the DM halo analysis in parts of the parameter space will not probably suffice to motivate the acquisition of a larger data set in this mode.  Compared with the choice of the CTA candidate array ($B$ or $E$) and the analysis method ({\\em Ring Method} or On-Off Methods), the WIMP annihilation spectrum has the strongest impact on the CTA sensitivity. The lower panel of Fig.~\\ref{fig:halo_methods_spectra} shows for both candidate arrays and the {\\em Ring Method} the limits obtained in the case of a dominant annihilation into $b\\bar{b}$ pairs (solid), $\\mu^+\\mu^-$ (dotted) and $\\tau^+\\tau^-$ (dashed). The small photon yield from $\\mu^+\\mu^-$ final states implies limits that are a factor of about 5 worse than limits for dominant annihilation into $\\tau^+\\tau^-$.  It is clear that the full potential of the halo analysis will be exploited by confronting individual DM models with their predicted WIMP annihilation spectra $dN_{\\gamma}/dE_{\\gamma}$ with data.    \\subsection{Anisotropies in the diffuse gamma-ray background\\label{sec:spatial}} Besides gamma-rays from individual resolved sources and Galactic foreground, another component of diffuse gamma-ray background radiation has been detected and proven to be nearly isotropic. This radiation dominantly originates from conventional unresolved point sources below the detection threshold, while another fraction might be generated by self-annihilating (or decaying) DM particles, which then could produce specific signatures in the anisotropy power spectrum of the diffuse gamma-ray background~\\cite{Ando:2007, Ando:2009, Cuoco:2008, Siegal-Gaskins:2008, Siegal-Gaskins:2010}. The different hypotheses about the origin of the gamma-ray background may be distinguishable by accurately measuring its anisotropy power spectrum.  Compared to the current generation of IACTs, CTA will have improved capabilities to measure anisotropies in the diffuse gamma-ray background, based upon a better angular resolution (determined by the point-spread function, PSF), an increased size of the FOV, and a higher background rejection efficiency. In the following, we discuss the effects of different assumptions on the background level and the anisotropy spectrum on the reconstruction of the power spectrum for the current generation of IACTs, and address the improvement obtainable with CTA. Finally, we make predictions for the discrimination between astrophysical and dark matter induced anisotropy power spectra for CTA.  \\subsection*{Simulation} In order to investigate the measured power spectrum and the impact of instrumental characteristics, a sample of \\emph{event lists} containing anisotropies generated with Monte-Carlo simulations was analyzed. The event lists were simulated by generating skymaps with a given anisotropy power spectrum. In total, $12$ skymaps covering the size of the FOV and being in different celestial positions were created, with a power spectrum for a given multipole moment $\\ell$ defined as $C_{\\ell} = 1/(2\\ell+1) \\sum \\lvert a_{\\ell m} \\rvert^2$, $m = -\\ell, \\dots, \\ell$, where $a_{\\ell m}$ denotes the coefficients of a (real-valued) spherical function decomposed into spherical harmonics. With $\\langle a_{\\ell m} \\rangle = 0$, $C_\\ell$ reflects the width of the $a_{\\ell m}$ distribution, which was assumed to be Gaussian. The simulations were made for different power spectra $\\ell (\\ell +1) C_{\\ell} \\sim \\ell^{s}$, with $s = 0.5,~1.0,~1.5,~2.0,~2.5$. The pixel size of these skymaps was $0.002^{\\circ}$, corresponding to $\\ell = 9 \\times 10^{4}$ (where $\\Theta_\\ell = 180^\\circ/\\ell$). The skymaps $I(\\vartheta, \\varphi)$ were normalized in a way that the pixel with the smallest signal was assigned the value~$0$ and the pixel with the largest signal was assigned $1$. Anisotropy power spectra were then derived from the fluctuation maps $I(\\vartheta, \\varphi)/\\langle I \\rangle$, such that for a \\textit{full signal} the maximum allowed difference in each map equals $1$. Note that this difference can be smaller when an additional isotropic \\textit{noise} component is present.  An event was simulated in three subsequent steps: First, the celestial position was randomly chosen within the FOV, and the event was classified to represent a \\textit{signal}- or isotropic \\textit{noise}-event, respectively. The decision for a signal event was based upon a normalized random number $z$: If $z$ was smaller than the skymap value at the corresponding position, the event was considered a signal event. Otherwise, another event position was selected while reapplying the procedure. Subsequently, the event map was convolved with a PSF of $0.1^\\circ$, which is similar to the resolution of current IACTs. The effect of a better angular resolution is discussed below. The event maps were simulated to contain $10^{7}$ entries. Note here that this number, as selected for the toy model, does in general not reflect the actual number of expected physical signal events. Therefore, the following discussion is focussed more on a qualitative discussion of the criticalities of the calculation rather than on making quantitative predictions.  To analyze an event list containing $N_{\\rm{ev}}$ events, a HEALPix skymap with $N_{\\rm{pix}}$ pixels was accordingly filled, and analyzed using the HEALPix software package\\footnote{http://healpix.jpl.nasa.gov/}. Therefore, the analyzed function is \\begin{equation} w(\\hat{n}) f(\\hat{n}) = \\frac{N_{\\text{pix}}}{N_{\\text{event}}} \\sum_{i = 1}^{N_{\\text{pix}}} x_{i} \\cdot b_{i}(\\hat{n}), \\end{equation} where $x_{i}$ denotes the number of events in pixel $i$, and $b_{i}(\\hat{n})$ equals $1$ inside pixel $i$ and $0$ outside. The function $w(\\hat{n})$ describes the windowing function --- in this case the FOV with Gaussian acceptance --- and $f(\\hat{n})$ denotes the original signal function over the full sky. The windowing function was normalized such that the integral over the full sky equals $4 \\pi$: \\begin{equation} \\label{eq_norm} \\int d\\Omega \\, w(\\hat{n}) f(\\hat{n}) = 4 \\pi. \\end{equation} Note that this differs from other analyses of this type, where $w$ is defined such that the maximum value is $1$. This difference in the normalization was done in order to keep a simple simulation code, and results should be equivalent. Final results were averaged over the corresponding $12$ skymaps.  \\subsubsection*{The effect of the anisotropy spectrum and the residual background level on the spectral reconstruction}  \\begin{figure}[h!t] \\centering \\includegraphics[width=0.95\\linewidth]{slope_Psf10.eps} \\caption{\\label{fig_Psf}Measured power spectra $l(l+1)C_\\ell/(2\\pi)$ for different slopes $s$ of the simulated input spectrum, compared to an isotropic background spectrum. Color-filled areas depict the RMS of the spectra. The size of the PSF is $\\sigma_{\\mathrm{PSF}} = 0.1^{\\circ}$. For reference, the simulated spectra are shown as dashed lines.} \\end{figure}  In Fig.~\\ref{fig_Psf}, we show the mean value and the RMS of the $C_{\\ell}$ power spectra. The value $C_{\\ell}$ represents the strength of anisotropies of the angular scale $\\Theta_{\\ell} = 180^{\\circ}/\\ell$. Anisotropies smaller than the angular resolution (defined by the PSF) are smeared out. This effect is clearly visible for large $\\ell\\geq1000$, where the power spectra converge into the Poissonian noise of the isotropic background spectrum. The angular resolution assumed for the simulation shown in this figure has a width of $\\sigma_{\\mathrm{PSF}} = 0.1^{\\circ}$. Furthermore, anisotropies with a size larger than the FOV are truncated at $\\ell\\sim100$ due to the effect of the windowing function. The simulated FOV in Fig.~\\ref{fig_Psf} has a width of $2.5^{\\circ}$, which is comparable with the FOV of current IACT experiments. For the toy model, Fig.~\\ref{fig_Psf} demonstrates that, for $\\ell \\sim 100\\!-\\!1000$, power spectra of different slopes are separable within the statistical errors and distinguishable from isotropic noise. CTA will have a smaller PSF as well as a larger FOV. This will make the signal vanish at larger $\\ell$ than in the example, and the windowing function will influence the spectrum to smaller $\\ell$ than in the figure. Therefore, we conclude that the FOV as well as the PSF, while important, will not be crucial for the investigation of anisotropies with CTA in the desired multipole range.  \\begin{figure}[h!t] \\centering \\includegraphics[width=0.95\\linewidth]{sn_slope05.eps} \\caption{\\label{fig_sn}Influence of the signal fraction on the measured power spectrum. Shown are the reproduced spectra for a slope $s = 0.5$ for several ratios between signal and total events; the \\textit{background} events are distributed isotropically. In order to estimate the effect of the noise ratio, the best fit levels are shown as dotted lines. The width of the PSF is chosen as in Fig.~\\ref{fig_reality}, $\\sigma_\\mathrm{PSF} = 0.05^\\circ$.} \\end{figure}  In general, the measured flux will be composed of both signal and background events. The background is produced mainly by two separate processes: \\begin{enumerate} \\item Events caused by cosmic rays (protons and electrons) which are misinterpreted as photon events. \\item An isotropic component of the photon background radiation, which does not count as \\textit{signal} according to our definition. \\end{enumerate} The influence of isotropic background is demonstrated in Fig.~\\ref{fig_sn}, where the power spectrum for $s = 0.5$ is shown for different background levels. Here, the signal fraction is defined by $f_\\mathrm{sig} = N_\\mathrm{sig}/N_\\mathrm{ev}$, where $N_\\mathrm{sig}$ denotes the number of signal events. The overall power is clearly reduced in case of fully isotropic background.  From the figure, we see that when the signal fraction improves by a factor 5, the power spectrum is boosted by about two orders of magnitude. For this reason, we expect the ten-fold improved CTA sensitivity to mark the major difference with respect to the current generation of IACTs for such studies.   \\subsubsection*{Prospects for astrophysical and dark matter anisotropies discrimination} \\begin{figure}[h!t] \\begin{center} \\includegraphics[width=0.95\\linewidth]{./compare_00300h_020.eps} \\end{center} \\caption{Comparison between measured anisotropy power spectra with (a) a pure astrophysical origin represented by $C_{\\ell} = 10^{-5}$ (blue bands) and (b) with an additional DM component, i.e., 20\\% of the total flux, represented by $C_{\\ell} = 10^{-3}$ (red bands). The assumed observation time is $300\\,\\mathrm{h}$. The three cases in each plot represent the hadronic background rates of $10\\,\\mathrm{Hz}$,$1\\,\\mathrm{Hz}$, and $0.1\\,\\mathrm{Hz}$.} \\label{fig_reality} \\end{figure}  The theoretical expectations for the power spectra of the diffuse gamma-ray flux of both the astrophysical as well as the DM components are highly model dependent. Since the astrophysical component is dominated by the gamma-ray flux from unresolved point sources, expected with a constant $C_{\\ell}$ ($s = 2.0$ in our notation), we conservatively assume the slope of the DM component(s) to be similar. In this scenario, the difference between the power spectra manifests in the normalization. For unresolved point sources, $C_{\\ell,\\,\\mathrm{blazars}} = 10^{-5}$, while for DM-induced anisotropies, considering the thermal annihilation cross-section $3\\times10^{-26}$~cm$^{3}$ s$^{-1}$, $C_{\\ell,\\,\\mathrm{DM}} = 10^{-3}$ is expected~\\cite[see, e.g.,][]{Siegal-Gaskins:2010}. In our simulation, this was realized by distributing $N=4 \\pi/C_{\\ell}$ point sources over the full sky. While representing a non-physical model, this is a convenient way of producing a Poissonian anisotropy power spectrum which is a reasonable assumption for generic astrophysical and DM emitters. The normalization of the signal was set by extrapolating the spectrum of the extragalactic gamma-ray background (EGB) \\cite{biB_Fermiiso} to $E > 100\\,\\mathrm{GeV}$. Note that the strength of the DM annihilation signal is strongly affected by the formation histories of DM halos and the distribution of DM subhalos. For example, Fig. 3 in \\cite{biB_Fermilimit} shows that the gamma-ray spectrum of DM annihilation could reach the measured gamma-ray background spectrum and therefore deliver a significant fraction of the measured flux. Here, we investigate the cases that $(a)$ the total EGB originates from astrophysical sources and $(b)$ 20\\% of the EGB (optimistically) originates from DM annihilation. The isotropic hadronic component depends on analysis cuts and the quality of the gamma-hadron separation. In the following, three different background rates are assumed: $10\\,\\mathrm{Hz}$, $1\\,\\mathrm{Hz}$, and an optimistic $0.1\\,\\mathrm{Hz}$ rate. We assume a CTA-like FOV with a radius of $5^\\circ$ and a CTA-like PSF with $\\sigma_\\mathrm{PSF} = 0.05^\\circ$. The results are shown in Fig. \\ref{fig_reality}, where each band represents a sample of 20 realizations. One can see in the figure that depending on the achieved background rate, in principle the two above mentioned scenarios $(a)$ and $(b)$ will be well distinguishable for CTA. ",
        "conclusions": "\\label{sec:conclusion}  In this study we have investigated the prospects for detection and characterization of several flavors of physics beyond the standard model with CTA.\\\\  \\subsection*{Particle Dark Matter searches} We have investigated dark matter (DM) searches with CTA for different observational strategies: from dwarf satellite galaxies (dSphs) in Section~\\ref{sec:dwarf}, from clusters of galaxies in Section~\\ref{sec:gc} and from the vicinity of the Galactic Centre in Section~\\ref{sec:halo}. In Section~\\ref{sec:spatial}, we discussed spatial signatures of DM in the diffuse extragalactic gamma-ray background.  Concerning searches in dSphs of the Milky Way, we have investigated the prospects for detection of well-known ``classical'' dSph like Ursa Minor and Sculptor, and one of the most promising ``ultra-faint'' dSph, Segue~1 (Table~\\ref{tab:jbar}). We have first shown that the predictions for core or cusp DM density profiles are quite similar for the baseline CTA angular resolution (Fig.~\\ref{fig:profiles}). We have then simulated a 100~h observation for several CTA arrays, and found that for Segue~1, we can exclude velocity-averaged cross-sections $\\left(\\sigma_{\\rm ann} v\\right)$ above $10^{-23}-10^{-24}\\,\\mbox{cm}^3\\mbox{s}^{-1}$ depending on different annihilation channels (Fig.~\\ref{fig:sigmav_ul}). We also presented the same results in terms of the minimum astrophysical factor for dSphs to be detected (Fig.~\\ref{jfactors}), showing that astrophysical factors of at least $10^{21}$~GeV$^{2}$~cm$^{-5}$ are needed. We finally showed the minimum intrinsic boost factor to achieve detection (Fig.~\\ref{fig:Boost}), which for Segue~1 is about 25 for a hard annihilation spectrum. The best candidate arrays for dSph study are array $B$ and $E$. Nevertheless, the robustness of our results is hindered by the yet not precise determination of the astrophysical factor in some cases.  Forthcoming detailed astronomical measurements will provide clues for deep exposure observations on the most promising dSphs, with, e.g., the planned SkyMapper Southern Sky Survey~\\cite{Keller:2007}, which will very likely provide the community with a new dSph population, complementing the Northern hemisphere population discovered by the SDSS. Also, the uncertainties on dark matter density will be significantly reduced by new measurements of individual stellar velocities available after the launch of the GAIA mission\\footnote{www.rssd.esa.int/Gaia}. Stacking-methods of Fermi-LAT dSphs data were proven valid to make constraints more stringent~\\cite{Abdo:2010b,GeringerSameth:2011iw, Cotta:2011pm}. The application of these methods for CTA is currently under study.  The search for DM signatures in galaxy clusters, investigated in Section~\\ref{sec:gc} was performed for two representative clusters, Perseus and Fornax. The former one is thought to have the highest CR-induced photon yield, and the latter is thought to have the strongest DM-induced signatures. Compared to dSphs, the gamma-ray signatures of galaxy clusters have several contributions: in the first place, the DM signal is expected from an extended region that can be larger than a few degrees, and secondly, gamma-rays induced by interactions of accelerated cosmic rays with the ambient fields and/or by individual cluster galaxies are an irreducible background to the DM signal, as recently shown in Refs.~\\cite{Aleksic:2010xk, Aleksic:Perseus}. We have simulated the prospects of detection in 100~h of observation by using MC simulations of extended sources. Regarding DM signatures, we have used the model of~\\citet{Pinzke:2011ek} for the Fornax cluster, and showed that in 100~h we could put contraints on the order of $\\left(\\sigma_{\\rm ann} v\\right)<10^{-25}\\,\\mbox{cm}^3\\mbox{s}^{-1}$ (Fig.~\\ref{fig:Fornax_DM}), which are competitive with respect to those obtained with dSphs. The results are promising: if the intrinsic boost factor from subhalos is larger than that predicted by the model we used, or mechanisms of Sommerfeld enhancement are at work, there is also the possibility to have a detection in $100-200$~h with array $B$ or $E$.  We have also considered the prospects of detection of CR-induced signal in hadronic acceleration scenarios in Fig.~\\ref{fig:Perseus_CR}. We have seen that the CR-induced emission from the Perseus cluster could be detected in about 100~h.  Finally, we discussed the more realistic case when DM-- and CR-induced gamma-rays are treated together. We discuss that the difference in both the spatial and spectral features of the two emissions can be used as a method for discrimination, while more quantitative results need dedicated MC which were not available when writing this contribution. We underline that the extension of the expected DM emitting region in galaxy clusters represents a problem for current Cherenkov Telescopes since their FOV is limited to $3-5$ degrees and their sensitivity rapidly decreases moving away from the centre of the camera. CTA will overcome this limitation, having a FOV of up to 10~deg and an almost flat sensitivity up to several degrees from the centre of the camera. For galaxy cluster searches, CTA will hence mark the difference compared to the current generation of IACTs.  More promising are DM searches of annihilation signatures in the Galactic halo, where the DM density is expected to be known with much higher precision than in the Galactic Centre itself or in (ultra-faint) dSphs or galaxy clusters. This was studied in Section~\\ref{sec:halo}. By adopting dedicated observational strategies of the region close to the Galactic Centre, as shown in Fig.~\\ref{fig:halo_rings}, it was shown that CTA has the potential to reach the thermal annihilation cross-section expected from WIMP DM of $10^{-26}\\,\\mbox{cm}^3\\mbox{s}^{-1}$ and lower (Fig.~\\ref{fig:halo_methods_spectra}) in 100~h observation of the vicinities the Galactic Centre using the ``Ring'' method. Models with a large photon yield from DM annihilation will be constrained for even smaller cross-sections. It is also expected that the limits presented here can be improved by factor of a few when the stereoscopic analysis of CTA events has been understood so well that a further suppression of the background becomes feasible. This would be the first time that ground-based Cherenkov telescopes could reach this sensitivity level.  Besides observations of individual dedicated objects, the capabilities of CTA for searching DM signals in the diffuse background of gamma-ray radiation were discussed in Section~\\ref{sec:spatial}. We discussed the reconstruction performance for different anisotropy power spectra and residual background level. Considering a current model for the anisotropy power spectra, we showed that CTA may be able to distinguish a DM-induced diffuse gamma-ray component from the astrophysical background.\\\\   \\begin{figure}[h!t] \\centering \\includegraphics[width=0.95\\linewidth]{./last_Plot.eps} \\caption{\\label{fig:Fermi.vs.CTA} Comparison of exclusion curves of Fermi-LAT in 24~months \\cite{Ackermann:2011} and expected for 10~years (rescaled with the square root of time). The exclusion curves for the various targets studied in this contribution are also reported for the $b\\bar{b}$ annihilation channel: for the dwarf satellite galaxy Segue~1 (green curve, see Sec.~\\ref{sec:dwarf}), for the Fornax galaxy cluster in case only DM-induced gamma-rays are considered (blue line, see Sec.~\\ref{sec:gc}) and for the ring-method of observation of the Galactic Centre vicinities (red line, see Sec.~\\ref{sec:halo}). } \\end{figure}   In Fig.~\\ref{fig:Fermi.vs.CTA}, we summarize the constraints that we expect with CTA for a WIMP annihilating purely into $b\\bar{b}$ in 100~h observation, with the different targets discussed above.  As already anticipated, the best results are expected for the observation of the vicinity of the Galactic Centre, where we expect to reach the thermal annihilation cross-section for WIMP DM of $10^{-26}\\,\\mbox{cm}^3\\mbox{s}^{-1}$.  Unlike present IACTs, whose sensitivity supersedes that of the Fermi-LAT at masses around a TeV, CTA will constitute the most sensitive instrument above masses of about 100~GeV. It should be noted that these estimates are conservative: the most important improvement can be expected from the possible redefinition and final optimization of the array layout\\footnote{For example, it is currently under disussion, the possibility to add 36 medium-size telescopes of Schwarzchild-Couder design to the arrays considered here. Preliminary simulations predict that improvement in the overall sensitivity by at least a factor of 2 compared to that studied here may be expected.}  In addition, the presented sensitivities were calculated using generic analysis which was not optimized specifically for the DM searches and thus our results could be considered conservative in this sense.  Obviously, a firm \\emph{identification} of DM requires a very good spectral discrimination with respect to any possible astrophysical background. Spectral shapes and (even more so) absolute normalization of these backgrounds are often poorly determined and the DM signal most likely is small in comparison. As the extent to which these factors affect detection claims is highly model and target dependent, we refer to future detailed more focused assessments.\\\\  It has been shown that the detection of gamma-rays provides complementary information to other experimental probes of particle DM, especially that of direct detection, because CTA could be able to access a fraction of the parameter space not accessible otherwise~\\cite{Bergstrom:2010gh, Bergstrom:2011}. With respect to particle searches at the LHC, the comparison is not straightforward, as LHC results are usually strongly related to specific models, and general conclusions are somewhat model dependent, as shown by recent publications from the ATLAS and CMS collaborations~\\cite{Khachatryan:2011tk, Aad:2011ks, Aad:2011hh}. Generically, the discovery of a candidate for particle DM will be limited by the available centre-of-mass energy. Other scenarios exist, in the context of specific super-symmetric models for DM, that exhibit parts of the model space not accessible by the LHC~\\cite{Profumo:2011zj}. In any case, LHC discovery of dark matter, would prompt the need for proof that the particle is actually consistent with the astrophysical DM, and close collaboration with LHC physicists is currently under organization to facilitate the optimal use of accelerator results within CTA. A concrete scenario has been analyzed by~\\citet{Bertone:2011pq} in the case of a SUSY model in the so-called co-annihilation region. Simulated LHC data were used to derive constraints on the particle physics nature of the DM, with the result that the LHC alone is not able to reconstruct the neutralino composition. The situation improves if the information from a detection of gamma-rays after the observation of the Draco dSph by CTA is added to the game: in this case the internal degeneracies of the SUSY parameter space are broken and including CTA allows us to fully interpret the particle detected at the LHC as the cosmological DM. In the other case where the LHC will not detect any physics beyond the Standard Model, predictions were made in the context of the CMSSM~\\cite{Bertone:2011kb} indicating that the mass of the neutralino will be bound to be larger than approximately $250$~GeV ($400$~GeV) if any new physics will be detected by the LHC for an energy of the centre-of-mass $\\sqrt{s}=14$~TeV and a luminosity of $1$~fb$^{-1}$ (100~fb$^{-1}$). In this scenario, CTA could be the only instrument to be able to detect and identify a WIMP candidate with masses beyond some hundreds GeV.   \\subsection*{Axion-like particle searches} In Section~\\ref{sec:axions}, the prospect of searches for axion-like particle (ALP) signatures with CTA were studied. We saw that the theoretical photon/ALP mixing has important implications for astronomical observations, in such that the mixing could distort the spectra of gamma-ray sources, such as Active Galactic Nuclei (AGN) (or galactic sources), in the TeV range. This distortion adds to that caused by the absorption of the gamma-ray photons with UV and IR photons of the Extragalactic Background Light (see Fig.\\ref{fig:sketch}).  The photon flux recently measured by some experiments, in particular at TeV energies, already exceeds that predicted by conventional models which attempt to explain spectra in terms of  observed source spectra and/ or EBL density~\\cite{aleksic11, aliu09, Nesphor1998, Stepanyan2002, Krennrich2008}, one should not expect a photon flux as high as recently measured by some experiments, in particular at TeV energies. The hard spectrum deduced for some AGN is difficult to explain with conventional physics as well. While it is still possible to solve these puzzles without exotic physics, photon/ALP conversions may naturally alleviate both problems. In order to quantitatively study the effect of photon-axion conversion over cosmological distances, the total photon flux from a simulated flare of a far-distant source was considered. The source was simulated based on the flat spectrum radio quasar 4C +21.35 (PKS 1222+21, $z = 0.432$, based on the observation performed by MAGIC~\\cite{pks1222_MAGIC}), assuming an intrinsic unbroken power-law spectrum, and trying to understand the observability under different absolute flux normalization and flare duration (Figs.~\\ref{fig:Ecrit200} and \\ref{fig:Ecrit1000}).  The range of characteristic scale energy (critical energy, $E_{crit}$) described in Eq.~(\\ref{eq:ecrit}), and thus the ALP mass that can be probed with CTA for the different ALP scenarios, is unknown and may be tested with CTA. In general, the distortion of the spectra due to ALP depends on the particular case, but as a general trend it will become larger as we increase the observation time and/or the flux of the source. As an example, we found that a 0.5 h duration flare like the one reported by MAGIC would not be enough for CTA to detect a significant effect in any of the tested ALP scenarios (Fig.~\\ref{fig:probabilities}). However, a flare of similar intensity, but lasting 5 h would already be enough to see the boost due to ALPs for those scenarios with E$_{crit} \\leq 500$ GeV. For a hypothetical flare with an intensity 5 times larger, lasting 5 hours, the accessible range of E$_{crit}$ would extend up to 1.3 TeV (Fig.~\\ref{fig:probabilities}). Hopefully, not only PKS 1222+21 but also many other similar objects will be followed-up by CTA in the near future, making the field of ALP searches very promising.  We must emphasize that a boost in the flux is {\\it only} possible in the energy range where the EBL is already at work. Thus, even for the most distant sources detected to date by IACTs, the energy range below $\\sim100$ GeV would not probably be of much help. On the other hand, even when $E_{crit}$ lies within the energy range covered by IACTs, the drop/jump might not be accessible to these instruments. This would be the case, for instance, if $E_{crit}$ is at the highest energies, from several to tens of TeV: the attenuation due to the EBL for a distant source would be huge, and the resulting flux, even after accounting for the ALP boost, too low to be detected by current IACTs or by CTA under any of the possible array configurations. Taking all these considerations into account, the most suitable energy for ALP searches with CTA seems to be an intermediate one in which the EBL is already present but still introduces only a moderate absorption, i.e. from a hundred GeV to a few TeV. As a result, we do not expect to obtain largely different results with candidate array configurations other than the one we used (array $E$), since they all perform very similarly in the intermediate energy range.  Finally, although very challenging given the uncertainty in the value of the IGMF, we should stress that the lack of detection of suspicious features in the spectra of distant gamma-ray sources might translate into useful constraints of the ALP parameter space (coupling constant and ALP mass). A more detailed study is definitely needed in order to find out what should be the best strategy to achieve the strongest constraints. This study will be done elsewhere.  \\subsection*{Lorentz Invariance Violation} In scenarios where Lorentz invariance is violated by quantum gravitational effects, the space-time fabric may be distorted so that the vacuum shows a non-unitary refractive index and thus the light speed would be wavelength dependent. Observation of gamma-ray flares from far distant objects like active galactic nuclei or gamma-ray bursts, may allow to detect the time--delay between photons of different energies not caused by intrinsic source mechanisms.  In Section~\\ref{sec:liv}, we discussed the sensitivity of the different CTA array configurations on detecting time--delays induced by Lorentz Invariance Violations (LIV). While limits on LIV from the current generation of IACTs are weaker than those estimated from Fermi-LAT measurement in the so-called linear case, CTA is likely to invert this scenario.  Using for the Maximum Likelihood Estimation method of \\citet{martinez09}, 500 Gaussian-shaped pulsed light-curves with time--delay from $-60$~s TeV$^{-1}$ to $60$~s TeV$^{-1}$ were simulated and reconstructed with the different arrays. CTA will have improved statistics of photons and larger spectral lever-arm due to the enlarged energy range with respect to the current generation of IACTs. This will allow to better differentiate between the two Gaussian peaks as shown in Fig.~\\ref{fig:jb1} for the different arrays. This ability to differentiate peaks is also discussed in Fig.~\\ref{fig:jb2} for different width of the peaks. The best array configurations were discussed. As a result of these studies, a gain of about a factor 50 in the LIV scale for the ``quadratic'' model (see Eq.~(\\ref{eq:liv})) is expected, compared to current generation of telescopes, while the limits on the linear term will largely exceed the Planck energy scale.  We then used extrapolation to high-energies of real AGN spectra observed by the current generation of IACTs for three representative scenarios: a very bright AGN (Mrk~421), a fast-variable one (PKS~2155-304) and a high-redshift one (3C~279). High-energy photons above 10~TeV will guarantee the best sensitivity to observe LIV signatures, and CTA with its improved sensitivity at those high energies, will allow to collect sufficient photons, whereas photon statistics will always be the final limiting factor on tests for time--delay.  Finally, pairs of realistic AGN lightcurves with and without time--delay were simulated, and folded with CTA performance. For each array, we calculated the fraction of photons in which the time--delay was successfully measured according to a quality factor $q$ (Eq.~\\ref{eq:qfactor}). In the most stringent case $(q<1)$, we report the photon fraction recovery of each individual array in Fig.~\\ref{fig:globalAGNHist}. As a main result, more than 10\\% of the time--delays can be recovered with several possible CTA arrays. If we relax the quality factor, and thus the precision on the reconstructed time--delay, essentially all the time--delays are recovered. We showed that arrays $C$, $D$, $H$, $I$ and $NB$, respectively for the Southern and Northern hemisphere, have the best chance for detection.   Based on these genuinely different time--delay reconstruction methods we ensure that our final results, with respect to CTA-array ranking, are free from any possible systematic effects related to a given analysis method e.g.\\ idealized source redshift-distribution, idealized source light-curves, and idealized source time-scales. In all analyses methods the sub-arrays $C$, $H$, $I$ and $NB$ seem to be sufficiently good to perform detection of LIV effects by measuring differences in the arrival times of VHE photons. That means that these arrays are in general sufficiently good to perform temporal studies of light-curve signals and even detection of time--delays in AGN induced intrinsically in the source. The latter is an interesting degeneracy, connected with the actual origin of the time--delays, that CTA will definitely be able to break through population studies not based on exceptional flaring states but on a routine basis.   \\subsection*{Other Searches} Finally, in Section~\\ref{sec:exotic}, we have qualitatively discussed the physics case of a selection other exotic physic searches which are in principle possible with CTA: the observation of atmospheric showers from $\\tau$-particles emerging from the Earth crust, the observation of atmospheric showers from magnetic monopoles, and the possible follow-up of gravitational waves events. Despite the prospects being sometimes pessimistic, those subjects were shown to underline again the possibility of using an astronomical observatory such as CTA for fundamental physics searches.\\\\   As a final closing remark, we believe that CTA could offer one of the most powerful tools in the study of some of the most pressing questions in modern physics. In the next few years it may lead to a range of new observables, new methods and new theories. In preparation for these developments, it is essential that work such as that performed here is continued.   \\subsection*{Acknowledgment} \\small{We gratefully acknowledge support from the following agencies and organisations: Ministerio de Ciencia, Tecnolog\\'ia e Innovaci\\'on Productiva (MinCyT), Comisi\\'on Nacional de Energ\\'ia At\\'omica (CNEA) and Consejo Nacional  de Investigaciones Cient\\'ificas y T\\'ecnicas (CONICET) Argentina; State Committee of Science of Armenia; Ministry for Research, CNRS-INSU and CNRS-IN2P3, Irfu-CEA, ANR, France; Max Planck Society, BMBF, DESY, Helmholtz Association, Germany; MIUR, Italy; Netherlands Research School for Astronomy (NOVA), Netherlands Organization for Scientific Research (NWO); Ministry of Science and Higher Education and the National Centre for Research and Development, Poland; MICINN support through the National R+D+I, CDTI funding plans and the CPAN and MultiDark Consolider-Ingenio 2010 programme, Spain; Swedish Research Council, Royal Swedish Academy of Sciences financed, Sweden; Swiss National Science Foundation (SNSF), Switzerland; Leverhulme Trust, Royal Society, Science and Technologies Facilities Council, Durham University, UK; National Science Foundation, Department of Energy, Argonne National Laboratory, University of California, University of Chicago, Iowa State University, Institute for Nuclear and Particle Astrophysics (INPAC-MRPI program), Washington University McDonnell Centre for the Space Sciences, USA.  We thank I.~Freire, S.~Funk, W.~Hofmann, A.~Murphy, A.~Pinzke, S.~Sarkar, D.~Torres and F.~Zandanel who provided comments on the manuscript.  D.~Emmanoulopoulos acknowledges the Science and Technology Facilities Council (STFC) for support under grant ST/G003084/1. This work was partially supported by the Spanish Consolider-Ingenio CPAN (CPAN09-PD13) and Multidark (CSD2009-00064). A. Jacholkowska and J. Bolmont acknowledge the support of GdR PCHE in France."
    },
    "1208/1208.5483_arXiv.txt": {
        "abstract": "{ Using precise full-sky observations from \\textit{Planck}, and applying several methods of component separation, we identify and characterize the emission from the Galactic ``haze'' at microwave wavelengths.  The  haze is a distinct component of diffuse Galactic emission, roughly centered on the Galactic centre, and extends to $|b|\\sim35\\degr$ in Galactic latitude and $|l|\\sim15\\degr$ in longitude.  By combining the \\textit{Planck} data with observations from the \\textit{Wilkinson Microwave Anisotropy Probe} we are able to determine the spectrum of this emission to high accuracy, unhindered by the large systematic biases present in previous analyses.  The derived spectrum is consistent with power-law emission with a spectral index of $-2.55 \\pm 0.05$, thus excluding free-free emission as the source  and instead favouring hard-spectrum synchrotron radiation from an electron population with a spectrum (number density per energy) $dN/dE \\propto E^{-2.1}$.  At Galactic latitudes $|b|<30\\degr$, the microwave haze morphology is consistent with that of the \\textit{Fermi} gamma-ray ``haze'' or ``bubbles,'' indicating that we have a multi-wavelength view of a distinct component of our Galaxy.  Given both the very hard spectrum and the extended nature of the emission, it is highly unlikely that the haze electrons result from supernova shocks in the Galactic disk. Instead, a new mechanism for cosmic-ray acceleration in the centre of our Galaxy is implied. } ",
        "introduction": "\\label{sec:introduction} The initial data release from the \\textit{Wilkinson Microwave Anisotropy Probe} (\\textit{WMAP}) revolutionised our understanding of both cosmology \\citep{spergel03} and the physical processes at work in the interstellar medium (ISM) of our own Galaxy \\citep{bennett03b}.  Some of the processes observed were expected, such as the thermal emission from dust grains, free-free emission (or thermal bremsstrahlung) from electron/ion scattering, and synchrotron emission due to shock-accelerated electrons interacting with the Galactic magnetic field.  Others, such as the anomalous microwave emission now identified as spinning dust emission from rapidly rotating tiny dust grains \\citep{draine98a,draine98b,deoliveiracosta02,finkbeiner04c,hinshaw07,boughn07,dobler08b,dobler09}, were more surprising. But perhaps most mysterious was a ``haze'' of emission discovered by \\citet{finkbeiner04a} that was centred on the Galactic centre (GC), appeared roughly spherically symmetric in profile, fell off roughly as the inverse distance from the GC, and was of unknown origin.  This haze was originally characterised as free-free emission by \\citet{finkbeiner04a} due to its apparently very hard spectrum, although it was not appreciated at the time how significant the systematic uncertainty in the measured spectrum was.  \\defcitealias{dobler08a}{DF08}  An analysis of the 3-year \\textit{WMAP} data by \\citet[hereafter DF08]{dobler08a} identified a source of systematic uncertainty in the determination of the haze spectrum that remains the key to determining the origin of the emission.  This uncertainty is due to residual foregrounds contaminating the cosmic microwave background (CMB) radiation estimate used in the analysis, and arises as a consequence of chance morphological correlations between the CMB and the haze itself. Nevertheless, the spectrum was found to be both significantly softer than free-free emission, and also significantly harder than the synchrotron emission observed elsewhere in the Galaxy as traced by the low-frequency synchrotron measurements of \\citet{haslam82} \\citep[see also][]{reich88,davies96,kogut07,strong11,kogut12}.  Finally, it was noted that this systematic uncertainty could be almost completely eliminated with data from the \\textit{Planck}\\footnote{\\Planck\\ (\\url{http://www.esa.int/Planck}) is a project of the European Space Agency (ESA) with instruments provided by two scientific consortia funded by ESA member states (in particular the lead countries France and Italy), with contributions from NASA (USA) and telescope reflectors provided by a collaboration between ESA and a scientific consortium led and funded by Denmark.} mission, which would produce estimates of the CMB signal that were significantly less contaminated by Galactic foregrounds.  The synchrotron nature of the microwave haze was substantially supported by the discovery of a gamma-ray counterpart to this emission by \\citet{dobler10} using data from the \\textit{Fermi Gamma-Ray Space Telescope}.  These observations were consistent with an inverse Compton (IC) signal generated by electrons with the same spectrum and amplitude as would yield the microwave haze at \\textit{WMAP} wavelengths.  Further work by \\citet{su10} showed that the \\textit{Fermi} haze appeared to have sharp edges and it was renamed the ``\\textit{Fermi} bubbles.'' Subsequently, there has been significant theoretical interest in determining the origin of the very hard spectrum of progenitor electrons.  Suggestions include enhanced supernova rates \\citep{biermann10}, a Galactic wind \\citep{crocker11}, a jet generated by accretion onto the central black hole \\citep{guo11a,guo11b}, and co-annihilation of dark matter (DM) particles in the Galactic halo \\citep{finkbeiner04b,hooper07,lin10,dobler11a}. However, while each of these scenarios can reproduce some of the properties of the haze/bubbles well, none can completely match all of the observed characteristics \\citep{dobler12}.  Moreover, despite the significant observational evidence, there have  been suggestions in the literature that the microwave haze is either an artefact of the analysis procedure \\citep{mertsch10} or not synchrotron emission \\citep{gold11}.  The former conclusion was initially supported by alternative analyses of the \\textit{WMAP} data that found no evidence of the haze \\citep{eriksen06, dickinson09}. However, more recently \\citet{pietrobon11} showed that these analyses, while extremely effective at cleaning the CMB of foregrounds and identifying likely contaminants of a known morphology (e.g., a low-level residual cosmological dipole), typically cannot separate the haze emission from a low-frequency combination of free-free, spinning dust, and softer synchrotron radiation.  The argument of \\citet{gold11} that the microwave haze is not synchrotron emission was based on the lack of detection of a polarised component.  This criticism was addressed by \\citet{dobler12} who showed that, even if the emission is not depolarised by turbulence in the magnetic field, such a polarised signal is not likely to be seen with \\textit{WMAP} given the noise in the data.  With the \\textit{Planck} data, we now have the ability not only to provide evidence for the existence of the microwave haze with an independent experiment, but also to eliminate the uncertainty in the spectrum of the emission which has hindered both observational and theoretical studies for nearly a decade.  In \\refsec{data} we describe the \\textit{Planck} data as well as some external templates we use in our analysis.  In \\refsec{methods} we describe the two most effective component separation techniques for studying the haze emission in temperature. In \\refsec{results} we discuss our results on the morphology and spectrum of the haze, before summarising in \\refsec{summary}. ",
        "conclusions": "\\label{sec:summary} We have identified the presence of a microwave haze in the \\textit{Planck} LFI data and performed a joint analysis with 7-year \\textit{WMAP} data.  Our findings verify not only that the haze is real, but also that it is consistent in amplitude and spectrum in these two different experiments.  Furthermore, we have used \\textit{Planck} HFI maps to generate a CMB estimate that is nearly completely clean of haze emission, implying that we have reduced systematic biases in the inferred spectrum to a negligible level.  We find that the unbiased haze spectrum is consistent with a power law of spectral index $\\beta_{\\rm H} = -2.55 \\pm 0.05$, ruling out free-free emission as a possible explanation, and strengthening the possibility of a hard synchrotron component origin. The spectrum of softer synchrotron emission found elsewhere in the Galaxy is $\\beta_{\\rm S} = -3.1$, consistent with a cosmic-ray electron population that has been accelerated in supernova shocks and diffused throughout the Galaxy. This spectrum is significantly softer than the haze emission, which is not consistent with supernova shock acceleration after taking into account energy losses from diffusion effects.  The microwave haze is detected in the \\planck\\ maps with both simple template regression against the data and a more sophisticated Gibbs sampling analysis.  The former provides an excellent visualisation of the haze at each wavelength on large scales while the latter allows a pixel-by-pixel analysis of the complete data set. While the template analysis allows us to derive the $\\beta_{\\rm H} = -2.55$ spectrum with high confidence, spectral determination with the Gibbs approach is more difficult given that noise must be added to the analysis to ensure convergence in the sampling method, and that a significantly more flexible model (in particular, one in which the spectrum of synchrotron is allowed to vary with each pixel) is used. However, not only is the spatial correspondence of the haze derived with the two methods excellent, but the Gibbs method allows us to show conclusively that the microwave haze is a separate component and not merely a variation in the spectral index of the synchrotron emission.  The morphology of the microwave haze is nearly identical from 23 to 44\\,GHz, implying that the spectrum does not vary significantly with position.  Although detection of the haze in polarisation with \\textit{WMAP} remains unlikely given the noise level of the data \\citep{dobler12}, future work with \\textit{Planck} will concentrate on using its enhanced sensitivity to search for this component."
    },
    "1208/1208.0868_arXiv.txt": {
        "abstract": "Spatially extended emission regions of Active Galactic Nuclei (AGN) respond to continuum variations, if such emission regions are powered by energy reprocessing of the continuum. The response from different parts of the reverberating region arrives at different times lagging behind the continuum variation. The lags can be used to map the geometry and kinematics of the emission region (i.e., reverberation mapping, RM). If the extended emission region is not spherically symmetric in configuration and velocity space, reverberation may produce astrometric offsets in the emission region photocenter as a function of time delay and velocity, detectable with future $\\mu$as to tens of $\\mu$as astrometry. Such astrometric responses provide independent constraints on the geometric and kinematic structure of the extended emission region, complementary to traditional reverberation mapping. In addition, astrometric RM is more sensitive to infer the inclination of a flattened geometry and the rotation angle of the extended emission region. ",
        "introduction": "\\label{sec:intro}  Reverberation mapping (RM) is a powerful tool to probe the structure of the broad emission line region (BLR) in AGNs and quasars \\citep[e.g.,][]{Blandford_McKee_1982,Peterson_1993} without the need to resolve the BLR. The basic ideas are that the BLR is powered by photoionization by the (ionizing) continuum from the accreting black hole and that it responds to the variations of the continuum in a light crossing time. The response from different parts of the BLR will arrive at different time delays with respect to the continuum variation. Thus by mapping a two-dimensional broad line response function in the time delay and velocity plane \\citep[velocity-delay maps, e.g.,][]{Blandford_McKee_1982,Horne_etal_2004} one can in principle recover the geometry and kinematics of the BLR. Over the past several decades, RM has proven to be a practical technique in studying the structure of BLRs. Although accurate velocity-time delay mapping is still lacking, RM studies have successfully measured average BLR sizes for several dozens of AGNs and quasars \\citep[e.g.,][]{Kaspi_etal_2000,Kaspi_etal_2007,Peterson_etal_2004,Bentz_etal_2009a,Bentz_etal_2009b}, and in a few cases, crude velocity-delay maps \\citep[e.g.,][]{Denney_etal_2009a,Bentz_etal_2010}.  RM studies have shown that the typical BLR size $R$ scales approximately with the AGN luminosity $L^{0.5}$. In the latest version of the $R-L$ relation \\citep[e.g.,][]{Bentz_etal_2009a}, $\\log (R/{\\rm light\\, days})=-21.3+0.519\\log(\\lambda L_{\\lambda}(5100\\textrm{\\AA})/{\\rm erg\\,s^{-1}})$. Thus for typical quasar luminosities ($\\lambda L_{\\lambda}(5100\\textrm{\\AA})=10^{45}\\,{\\rm erg\\,s^{-1}}$, or bolometric luminosity $L_{\\rm bol}\\approx 10^{46}\\,{\\rm erg\\,s^{-1}}$) at $z\\sim 0.5$, the BLR size is $\\sim 0.1\\,{\\rm pc}$ ($\\sim 15\\,\\mu$as). This scale is 3 orders of magnitude smaller than the diffraction limit of 10m class optical telescopes (${\\cal \\theta}\\sim$ tens of mas). Optical/near-IR interferometry with $\\mu$as resolution has yet to come. Thus reverberation mapping will remain one of the few practical methods to measure the BLR size, along with microlensing in gravitationally lensed quasars \\citep[e.g.,][]{Morgan_etal_2010, Sluse_etal_2011}.  However, measuring the source photocenter positions can achieve a factor of $\\sim 1/\\sqrt{N_{\\rm photon}}$ enhancement in precision compared with the image resolution \\citep[e.g.,][]{Lindegren_1978,Bailey_98a}, where $N_{\\rm photon}$ is the number of photons received in a bandpass. This means for $10^6$ photons, the achievable astrometry precision is 3 orders of magnitude smaller than the image resolution. One application of this idea is {\\em spectroastrometry} \\citep[e.g.,][]{Bailey_98a}, where differential astrometric positions as a function of wavelength (velocity) can be used to probe otherwise unresolved sources \\citep[e.g.,][]{Bailey_98b,Gnerucci_etal_2011}.  When the astrometric precision approaches the angular BLR size, it becomes possible to study the BLR structure with astrometric signatures. A simple application would be to use spectroastrometry to place constraints on the BLR structure\\footnote{For instance, in certain geometries of the BLR, the blue and red parts of the broad line emission may have offset photocenters. The detection of such photocenter offsets can put constraints on the BLR size and kinematics. }. A more radical possibility, however, is to detect and model the wobble in the broad line emission photocenter due to reverberation to the continuum variations, since the different arrival times of the BLR response will cause shifts in the observed BLR photocenter.  In this work we investigate the feasibility of combing the traditional intensity RM with astrometric information. Unlike spectroastrometry, the reverberation of the BLR will directly induce shifts in the photocenter position of broad line emission as a function of time delay and velocity. As we will show below, the pattern of photocenter shifts is determined by the geometry and kinematics of the BLR, thus providing independent constraints on the BLR structure, complementary to traditional intensity RM. Although our main focus is on the BLR, this method can be readily applied to the reverberation of the dust torus of AGNs \\citep[e.g.,][]{Suganuma_etal_2006} or other reverberation systems, where the required astrometric precision might be less stringent. We describe the basic formalism in \\S\\ref{sec:form} and demonstrate this method in \\S\\ref{sec:model} and \\S\\ref{sec:result} with simple models for the BLR. We briefly discuss the practical issues with this method in \\S\\ref{sec:disc}, with an overview on the perspectives of achieving the required astrometric precision in the near future in \\S\\ref{sec:disc_astro}, and conclude in \\S\\ref{sec:con}. ",
        "conclusions": "\\label{sec:con}  In this paper we outlined a simple idea of astrometric reverberation mapping: observing the astrometric ``wobble'' of the BLR photocenter in reverberation events. We demonstrated its potential in constraining the geometry and kinematic structure of the BLR with simple BLR models.  The virtue of astrometric RM is that it provides an independent set of time series that can be combined with the intensity RM data to constrain the BLR model. In addition, it is more sensitive to the inclination of a disk-like geometry\\footnote{Other methods to infer the inclination of a disk geometry of the BLR include polarimetry \\citep[e.g.,][]{Li_etal_2009} and double-peaked broad line profile fitting \\citep[e.g.,][]{Eracleous_Halpern_1994}.} and inference of its angular momentum direction, than intensity RM alone. Although the BLR is still unresolved, the spatial information from precision astrometry for the reverberating parts of the BLR greatly facilitates using reverberation mapping to probe the BLR structure. Once we have a set of high quality continuum and broad line flux monitoring time series with precision photometry and astrometry, we can fit the data in the time domain using model intensity light curves and photocenter evolution. General Markov Chain Monte Carlo approaches with Bayesian inference \\citep[e.g.,][]{Pancoast_etal_2011,Brewer_etal_2011} can be applied to search the parameter space of different model families and to quantify model uncertainties.  While conceptually simple, this method is not yet ready for immediate applications. We discussed realistic timescales for achieving the required astrometric precision ($\\mu$as to tens of $\\mu$as), and concluded that ground-based large-aperture telescopes with AO offer the best hope to apply this method to luminous quasars and dust torus RM within a decade or two, and a space-based interferometer (a SIM/SIM-Lite like mission) with $\\mu$as astrometric precision in the next decade or two may extend this method to the Seyfert regime."
    },
    "1208/1208.5518_arXiv.txt": {
        "abstract": "We have used the Caltech Submillimeter Observatory (CSO) to follow-up a sample of WISE-selected, hyperluminous galaxies, so called W1W2-dropout galaxies. This is a rare ($\\sim 1000$ all-sky) population of galaxies at high redshift (peaks at $z$=2-3), that are faint or undetected by WISE at 3.4 and 4.6 $\\mu$m, yet are clearly detected at 12 and 22 $\\mu$m. The optical spectra of most of these galaxies show significant AGN activity. We observed 14 high-redshift ($z > 1.7$) W1W2-dropout galaxies with SHARC-II at 350 to 850 $\\mu$m, with 9 detections; and observed 18 with Bolocam at 1.1 mm, with five detections. Warm {\\it Spitzer} follow-up of 25 targets at 3.6 and 4.5 $\\mu m$, as well as optical spectra of 12 targets are also presented in the paper. Combining WISE data with observations from warm {\\it Spitzer} and CSO, we constructed their mid-IR to millimeter spectral energy distributions (SEDs). These SEDs have a consistent shape, showing significantly higher mid-IR to submm ratios than other galaxy templates, suggesting a hotter dust temperature. We estimate their dust temperatures to be $60-120$ K using a single-temperature model. Their infrared luminosities are well over 10$^{13}$\\,L$_\\odot$. These SEDs are not well fitted with existing galaxy templates, suggesting they are a new population with very high luminosity and hot dust. They are likely among the most luminous galaxies in the Universe. We argue that they are extreme cases of luminous, hot dust-obscured galaxies (DOGs), possibly representing a short evolutionary phase during galaxy merging and evolution. A better understanding of their long-wavelength properties needs ALMA as well as {\\it Herschel} data. ",
        "introduction": "\\label{intro}  The redshift $z\\sim$2-3 epoch stands out as a unique era for studying galaxy formation and evolution. At this epoch, the cosmic star formation rate reaches its peak (Heavens et al. 2004; Hopkins \\& Beacom 2006;  Reddy et al. 2008), and ultra-luminous infrared galaxies (ULIRGs, $L_{IR} > 10^{12} L_{\\odot}$, Sanders \\& Mirabel 1996) contribute a significant fraction to the infrared luminosity density (Elbaz et al. 2002,  Chapman et al. 2005, Caputi et al. 2007, Reddy et al. 2008, Magnelli et al. 2009). The cosmic quasar density also peaks around $z\\sim$2 (Hopkins et al. 2007, Assef et al. 2011). A framework of galaxy evolution through major mergers has been gradually built up by theorists (Barnes \\& Hernquist 1992, Schweizer 1998, Jogee 2006, Hopkins et al. 2006, 2008). In one of the most popular scenarios (e.g. Hopkins et al. 2008), the tidal torques generated by major mergers funnel gas into the center of galaxies, leading to a central starburst and rapid growth of a supermassive black hole (SMBH). Black hole and supernova feedback terminate further star formation, evacuating the residual gas and dust, leaving a visible quasar and remnant spheroid. This picture establishes the evolutionary connections between ULIRGs, quasars, and massive elliptical galaxies.  Submillimeter galaxies (SMGs) are thought to be the analogues of local ULIRGs at high redshift (Blain et al. 2002, Tacconi et al. 2008). SMGs are selected by their strong cold dust emission at 850 $\\mu$m ($F_{850} > $5 ~mJy). They are characterized by very high star formation rates (100-1000 M$_{\\odot}$yr$^{-1}$) and infrared luminosity ($L_{IR}\\sim 8\\times 10^{12} L_{\\odot}$, Chapman et al. 2005, Magnelli et al 2012). Although most SMGs host growing black holes (e.g., Alexander et al. 2005, 2008), their luminosities are normally dominated by star formation (Swinbank et al. 2004; Men\\'{e}ndez-Delmestre et al. 2007, Younger et al. 2008, Hainline et al. 2011). The redshift distribution of SMGs strongly peaks at $z$=2-3 (Chapman et al. 2005), and the surface density of SMGs is several hundred per square degree.  An 850$\\mu$m selected sample (SMGs) may be biased toward ULIRGs with large amounts of dust, but miss a substantial population of ULIRGs with a smaller amount of (but warmer) dust, which can be found by surveys at shorter wavelengths. A series of surveys using bright {\\it Spitzer} 24$\\mu$m emission combined with optically faint photometry have been carried out to probe the ULIRG population with emission from smaller and warmer dust grains (e.g., Rigby et al. 2004, Donley et al. 2007, Yan et al. 2007, Farrah et al. 2008, Soifer et al. 2008, Lonsdale et al. 2009, Huang et al. 2009). One of the simplest search criteria is given as $F_{24} > 0.3$ mJy, and $R-[24]> 14$ (where $R$ and [24] are Vega magnitudes for $R$ band and {\\it Spitzer} 24 $\\mu$m), or roughly $F_{24}/F_{R} > 1000$ (Dey et al. 2008, Fiore et al. 2008), leading to a well defined $z \\sim$2 population which is referred to as Dust Obscured Galaxies (Dey et al. 2008, hereafter D08). The most luminous DOGs have star formation rates (500-1000 M$_{\\odot}$yr$^{-1}$ or more) and infrared luminosities ($L_{IR}\\sim 10^{13} L_{\\odot}$, Bussmann et al 2009, Tyler et al. 2009, Melbourne et al. 2012) that are comparable to SMGs. It has been proposed that both SMGs and DOGs are an early phase of galaxy merging, with SMGs representing an earlier, starburst-dominant phase, while luminous DOGs are in a transitional phase from starburst-dominated to AGN-dominated (e.g. Narayanan 2010). The bolometric luminosities also reach their maximum during these phases, making the most luminous galaxies in these phases also among the most luminous objects in the Universe.  Looking for the most luminous galaxies in the Universe is one of the major goals of NASA's Wide-field Infrared Survey Explorer (WISE, Wright et al. 2010). WISE surveyed the entire sky at 3.4, 4.6, 12 and 22 $\\mu$m (hereafter W1, W2, W3, W4) in 2010. The WISE dataset is well suited to investigate the starburst-AGN phase of distant, infrared luminous galaxies. At $z \\sim$ 2-3, starburst- or AGN-heated hot dust can be traced by 12 and 22 $\\mu$m emission, while the rest near infrared (NIR) obscuration is sampled by 3.4 and 4.6 $\\mu$m continuum. Studies of luminous infrared galaxies with the WISE W1, W2 and W4 bands can take advantage of existing knowledge and techniques developed by earlier studies with {\\it Spitzer} at similar wavelengths (IRAC at 3.6 and 4.5 $\\mu$m, and MIPS at 24 $\\mu$m). Observing W4 selected galaxies with WISE is similar to observing 24 $\\mu$m bright galaxies with {\\it Spitzer}, but with the surveyed area enlarged from a few tens of square degrees covered by existing DOG surveys to the entire sky.  In order to search for hyperluminous infrared galaxies (HyLIRGs, $L_{IR} > 10^{13} L_{\\odot}$) from the WISE dataset, the WISE team has explored multiple methods to select candidates. The most productive method so far has been to search for more heavily obscured galaxies, whose W1 (3.4 $\\mu$m) and W2 (4.6 $\\mu$m) flux densities become faint or undetected by WISE, while remaining easily detectable at 12 and/or 22 $\\mu$m, with typical W4 (22 $\\mu$m) flux densities $> 7 {\\rm mJy}$. We call this population \"W1W2-dropouts\" (Eisenhardt et al 2012) or for brevity \"W12drops\".  Follow-up spectroscopy of more than 100 W12drop galaxies at large telescopes (this paper, Eisenhardt et al. (2012), see also Bridge et al. 2012) reveals that a large fraction ($> 65\\%$) of these galaxies are at high redshift ($z > 1.5$), with the highest at $z$=4.6. Most of the redshifts are between 2 and 3, suggesting they also trace the peak epoch of cosmic star formation and QSO activity.  At these redshifts, such high flux densities at 22 $\\mu$m imply extremely high luminosities. They are potentially hyperluminous galaxies. In order to understand the dust properties and calculate the total luminosities of these unusual galaxies, continuum measurements at longer wavelengths are crucial. As the first high redshift examples were identified, we began follow-up 0.35-1.1 mm continuum observations using the Caltech Submillimeter Observatory (CSO), in order to construct their SEDs and to explore the nature of W12drop galaxies.  In this paper, we report the initial results of this follow-up study. The WISE data are described in section 2.1, and the W12drop population followed up with the CSO and reported here is listed in Table~1.  Sections 2.2 and 2.3 describe the CSO data, while section 2.4 describes the follow-up optical spectroscopy, which is summarized in Table 2. Section 2.5 describes Spitzer follow-up observations at 3.6 $\\mu$m and 4.5 $\\mu$m of the W12drops, which were selected to be difficult to detect by WISE at W1 (3.4 $\\mu$m) and W2 (4.6 $\\mu$m), and the photometry for the sources is presented in Table 3.  Section 3 presents luminosity and dust temperatures constraints from the photometry, while section 4 compares W12drop properties to those of DOGs and SMGs, and section 5 summarizes the findings. Throughout this paper we assume a \u039bCDM cosmology with  $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.27$, and $\\Omega_{\\Lambda} = 0.73$ . ",
        "conclusions": "\\label{sourcedisc}  Our CSO follow-up observations of 26 W12drop galaxies show that their luminosities are very high, with some over $10^{14} L_{\\odot}$, and a median and mean of 5.7 and 6.1 $\\times 10^{13} L_{\\odot}$, all using the conservative power-law method. This is roughly an order of magnitude brighter than the typical SMG (with median luminosity $L\\sim 8\\times 10^{12} L_{\\odot}$,  Chapman et al. 2005, Kov\\'{a}cs et al. 2006), or DOG (with median and mean luminosity $\\sim 6 \\times 10^{12}$  and $9 \\times 10^{12} L_{\\odot}$, Melbourne et al. 2012), and is comparable to the brightest known optically selected quasars (Schneider et al. 2005). Extremely luminous infrared galaxies are often found to be magnified by galaxy-galaxy lensing (Blain 1996, Eisenhardt et al. 1996, Solomon \\& Vanden Bout 2005, Vieira et al. 2010, Negrello et al. 2010).  An immediate concern about the hyperluminous W12drop galaxies is whether they are lensed, too. However, high-resolution imaging follow-up of selected W12drops does not indicate gravitational lensing (Eisenhardt et al. 2012, Bridge et al. in prep.), so that the derived $\\sim 10^{14} L_{\\odot}$ luminosities are consistent with being intrinsic based on the current data. Additional high-resolution follow-up observations with Hubble Space Telescope are currently underway, and should reveal if these W12drop galaxies are not lensed. If the lack of lensing is confirmed, these galaxies are one of the most luminous populations in the Universe.  Their unusually high dust temperatures and extremely high luminosities make W12drop galaxies of great interest for studying galaxy formation and evolution. How do they become so luminous? Are they experiencing special evolutionary events? What is their relationship to other well-established galaxy populations, such as SMGs and DOGs?  Classical SMGs are defined with strong 850$\\mu$m emission ($>$ 5mJy) which normally indicates significant cold-dust content. Table 3 gives examples of some W12drop galaxies that meet this criterion. Hence some W12drop galaxies would be selected as SMGs. However, the relatively low detection rate with Bolocam at 1.1 mm implies that many W12drop galaxies are not as bright as SMGs at longer wavelengths. This is understandable given that W12drop galaxies are dominated by emission from hotter dust. DOGs normally have both AGN and starburst contributions, with warmer dust grains than SMGs. In Dey et al (2008), DOGs are defined as galaxies with $F_{24} > 0.3$ mJy, and $R-[24]> 14$ (in Vega magnitudes), where the $R$ photometry is centered at 6393 \\AA. Since the W4 band at 22 $\\mu$m is similar to the {\\it Spitzer} 24 $\\mu$m band, our W4 $>$ 7 mJy selection corresponds to much higher flux densities at 24 $\\mu$m than normal DOGs. To make a comparison between W12drops and typical DOGs, we obtained $r$-band (centered at 6231 $\\AA$) photometry from SDSS (DR8), as listed in table 3, and used the $r$-W4 color to approximate the $R$-[24] color. Taking the average power-law index $\\alpha$ of 2.09 from Table 4 and extrapolating to $r$-band and 24 $\\mu$m, the difference between the $r$-W4 and $R$-[24] color ranges from 0.2 to 0.24 mag as $R$-[24] changes from 14 to 17. All 18 targets in Table 3 that are covered by SDSS DR8 meet the $r$-W4 $>$ 14 DOG criterion, with $r$-W4 ranging from 14.4 to 16.1 for $r$-band detected sources, and $r$-W4 $>$ 15.3 for $r$-band undetected sources (using $r$=22.9 Vega mag as the SDSS detection limit). In figure 6 we compare the distribution of $R$ vs. $R$-[24]  for these high-redshift W12drops to DOGs in D08. Clearly all W12drop galaxies in Table 3 qualify as DOGs, with similar colors, but are much brighter at 24 $\\mu$m.  Although many and maybe most W12drop galaxies can be classified as DOGs, their properties are quite different from normal DOGs. Comparing to the DOGs reported in D08, W12drop galaxies have an order of magnitude higher luminosity, although their redshift distributions are similar (Eisenhardt et al. 2012). Bussmann et al. (2009) used SHARC-II at 350 $\\mu$m to follow-up a subset of DOGs with the brightest 24 $\\mu$m flux densities from D08 (Fig 6), obtaining infrared luminosities (8-1000$\\mu$m, $\\sim 10^{13} L_{\\odot}$) and dust temperatures ($>$30-60 K), still significantly lower than for the W12drop galaxies reported here. Since the D08 survey covered only $\\sim $ 9 deg$^{2}$ (Bo$\\ddot{\\rm o}$tes field), the DOG surface density is $\\sim$ 320 DOGs per square degree with 24 $\\mu$m fluxes density greater than 0.3 mJy. The W12drop selection requires W4 $>$ 7 mJy, which is at the high 24 $\\mu$m flux density end of the D08 sample, and only selects $\\sim$ 1000 targets over the whole sky. The typical DOGs are 20 times fainter than the W12drops, but the latter are about 10000 times rarer. The high-z W12drop galaxies are apparently extreme cases of DOGs with very high dust obscuration and hotter dust temperatures, and appear to be hyperluminous, hot DOGs.  The very low surface density of W12drop galaxies suggests either they are intrinsically extremely rare, or are only seen during a very short phase of galaxy evolution. DOGs are thought to be the transitional phase of mergers between starburst-dominated mode and AGN-dominated modes. Based on the mid-IR SED (3.6-24 $\\mu$m), D08 classified DOGs into two categories, those which have a distinct ``bump'' in their SED between 3-10 $\\mu$m attributed to the redshifted starlight from rest-frame 1.6 $\\mu$m, and those whose mid-IR SED is a power-law. Bump DOGs are thought to be dominated by starbursts (Yan et al. 2005, Sajina et al. 2007, Farrah et al. 2008, Desai et al. 2009), and tend to have fainter 24 $\\mu$m flux densities (Dey et al. 2008), while power-law DOGs are thought to be dominated by AGN in the mid-IR (Weedman et al. 2006, Donley et al. 2007, Yan et al. 2007, Murphy et al. 2009), and make up most of the bright end of the 24$\\mu$m flux density distribution.  The fraction of power-law DOGs increases from 10\\% at $F_{24 \\mu m}=0.3$ mJy to 60\\% at $F_{24 \\mu \\rm m}= 1$ mJy in the {\\it Spitzer} Deep, Wide-Field Survey (Ashby et al. 2009). The mid-IR (MIPS 24 $\\mu m$) to submm (SPIRE 250 $\\mu m$) flux density ratio for power-law DOGs is found to be similar to the AGN dominated ULIRG Mrk231 (Melbourne et al. 2012). The IRAC1 to W4 SEDs of W12drops (Fig. 5) are more like the mid-IR SEDs of DOGs rather than SMGs' (Hainline et al. 2009). They show typical power-law shapes with no obvious bumps, and are very bright at 24 $\\mu$m. Consequently, it is plausible that W12drop galaxies are also dominated by very powerful AGNs. These powerful, highly obscured AGNs can heat the surrounding dust cocoon to a very high temperature.  Although the SEDs of W12drops are dominated by emission from very hot dust components (Figure 5) that are likely contributed by powerful AGNs, a hot AGN component alone can't explain all of the observed SED from the mid-IR to millimeter bands. It is likely that the SEDs are composed of multiple components with different temperatures. A more detailed model to decompose SEDs with multiple-temperature components needs a more complete set of SED data, which will become available from our ongoing {\\it Herschel} program. But the 350$\\mu$m and 1.1 mm data reported in this paper can give a useful constraint on the coldest component, if we assume that the 350 $\\mu$m to 1.1 mm SED is tracing the coldest dust in these galaxies. In figure 7, we plot the modeled flux density ratios of 350 $\\mu$m to 1.1 mm continuum versus the redshifts, for models with a single-temperature black-body times a wavelength-dependent opacity, with various dust temperatures and emissivities. W12drop galaxies with available 350 $\\mu$m and 1.1mm measurements are plotted in the figure. For comparison, models based on galaxies with significant starburst components (Arp220 and Mrk231) are also plotted. For $\\beta$ between 1.5 and 2.0, the 350 $\\mu$m to 1.1 mm ratios of W12drop galaxies in Figure 7 apparently favor a model with T$_{\\rm dust}$ less than 50 K, in addition to the $\\sim$ 100 K hot dust component that dominates the mid-IR.  This temperature of colder dust is comparable to the typical dust temperature of $\\sim$35 K associated with starburst galaxies (such as Arp220). Considering that the very hot AGN component will contribute more to the continuum at 350$\\mu$m than at 1.1mm, the actual 350 $\\mu$m to 1.1 mm ratios that trace the coldest dust could be lower, and therefore closer to the track of Arp220 or Mrk231 in Figure 7. This may imply that the cold dust component in these galaxies is not very different from those in starburst galaxies, and is possibly related to star formation. For example, in the detailed study of the first discovered W12drop galaxy (W1814+3412, Eisenhardt et al. 2012), a significant starburst is found, although only contributing a small fraction to the overall luminosity. The cold dust properties of W12drop galaxies may also be different from known obscured QSOs (e.g. Mart\\'{i}nez-Sansigre et al. 2009). The ratio of 350 $\\mu$m to 1.1 mm in AMS16, a high-z obscured quasar, is lower than for w12drops, as plotted in Figure 7. A detailed study of the long-wavelength properties for these W12drop galaxies, for instance, to distinguish the contribution and distribution of cold dust (star formation) and hot dust (AGN), will need observations from ALMA as well as {\\it Herschel}.  The similarity between the optical to 22 $\\mu$m SEDs of DOGs and W12drop galaxies, with the latter being much brighter, suggests W12drops may be the high luminosity tail of the DOG distribution. But the high mid-IR to submm luminosity ratio of W12drops implies they are much hotter than typical DOGs. Are these W12drops merely luminous DOGs, or a distinct population? Or do they have any evolutionary connection? Some theoretical models for DOGs (e.g. Narayanan et al. 2010) propose that SMGs, bump DOGs, and power-law DOGs may form an evolutionary sequence, representing the transition of merging galaxies from a starburst-dominated phase to an AGN-dominated phase, although direct observational support for this is still rare. In Figure 4a, we see a strong correlation between the mid-IR flux density and the mid-IR to submm luminosity ratio that supports such a sequence, with W12drops at the highest luminosities. The correlation is roughly linear, suggesting the cold dust component (traced by 350$\\mu m$ emission which may be from a starburst) doesn't change significantly during this process, as clearly shown in Figure 4b, while the hot dust component (traced by 24$\\mu m$) becomes stronger, possibly tracing the growth of an embedded SMBH. In this scenario, W12drops represent a late phase of this evolution, with more massive SMBHs and similar cold-dust components to SMGs and DOGs. If so, the low surface density of W12drops suggests either such a phase is very short, or not every galaxy goes through this stage. A better understanding of whether the populations have an evolutionary connection will need a fuller study of the W12drop population, and of the luminosity function of all these populations."
    },
    "1208/1208.2424_arXiv.txt": {
        "abstract": "We present results based on Chandra observations of a large sample of 129 hot galaxy clusters. We measure the concentration parameter c$_{200}$, the dark mass M$_{200}$ and the baryonic mass content in all the objects of our sample, providing the largest dataset of mass parameters for galaxy clusters in the redshift range $z$ = 0.01 -- 1.4. We confirm that a tight correlation between c$_{200}$ and M$_{200}$, $c \\propto M^a_{vir}/(1+z)^b$ with $a = $ -0.56 $\\pm$ 0.15 and $b = $0.80$ \\pm$ 0.25 (68 per cent confidence limits), is present, in good agreement with the predictions from numerical simulations and previous observations. Fitting the mass profile with a generalized NFW model, we got the inner slope $\\alpha$, with $\\alpha = 0.94 \\pm 0.13$. Finally, we show that the inner slope of the density profile, $\\alpha$ correlates with the baryonic mass content, $M_b$: namely  $\\alpha$ is decreasing with increasing baryonic mass content. ",
        "introduction": "\\label{sec:1}  \\indent\\indent  The problem of dark matter halos formation has a long history going back to the seminal paper of \\citet{Gunn:72} and \\citet{Gunn:77}, who studied the density profile formation using the collapse of a spherical perturbation in an expanding background. The quoted papers and several following analytical models (e.g., \\citealt{Fillmore:84}; \\citealt{Bertschinger:85}; \\citealt{Hoffman:85}), found that density profiles are described by power-laws in all the radius range. More recent semi-analytical models (e.g., \\citealt{White:92}, \\citealt{Subramanian:00}; \\citealt{Delpopolo:00}; \\citealt{El-Zant:01}, \\citealt{El-Zant:04}; \\citealt{Hiotelis:02}; \\citealt{LeDelliou:03}; \\citealt{Ascasibar:04}; \\citealt{Williams:04}; \\citealt{Tonini:06}; \\citealt{Ascasibar:07}) showed that the profile is not a power-law, similarly to N-body simulations results (e.g., \\citet{Navarro:96}, \\citet{Navarro:97}); Navarro--Frenk--White (NFW)), \\citealt{Moore:98}; \\citealt{Jing:00}; \\citealt{Klypin:01}, \\citealt{Bullock:01}, \\citealt{Power:03}, and \\citealt{Navarro:04}, \\citealt{Navarro:10}) found that the spherically averaged density profiles of the N-body DM halos are similar, regardless of the mass of the halo or the cosmological model. The NFW profile is given by: \\begin{equation} \\rho(r)= \\frac{\\rho_0}{r/r_s(1+r/r_s)^2}=\\frac{\\rho_{critic} \\Delta_c }{r/r_s(1+r/r_s)^2} \\end{equation} where $\\rho_{critic}$ is the critical density of the Universe at the cluster's redshift $z$, and $\\Delta_c $ is the virial overdensity. The scale radius $r_s$ is connected to the virial radius $r_{vir}$ through the concentration parameter $c$, $c= r_{vir}/r_s$\\footnote{Since $r_{vir}$ is difficult to determine observationally, its value is often approximated by the radius within which the average density is greater than the critical density by a specified factor (e.g., 200). In the following of the paper, $r_{vir}$ is identified with $r_{200}$.}. At small scales, the logarithmic density slope, that is also known as inner slope, is given by \\begin{equation} \\alpha= -\\frac{d \\log{\\rho}}{d  \\log{r}}|_{r \\rightarrow 0}=1 \\end{equation} The quoted profile diverges as $\\rho \\propto r^{-1}$ in the inner part, and at large radii behaves as $\\rho \\propto r^{-3}$. The inner slope is even steeper in \\citet{Moore:98} profile ($\\rho \\propto r^{-1.5}$). More recent simulations (e.g., \\citealt{Power:03}, \\citealt{Hayashi:04}, and \\citealt{Navarro:04}, \\citealt{Navarro:10}; \\citealt{Stadel:09}) showed that density profiles are better fitted by the Einasto profile, which becomes shallower towards the centre of the halo.  Unfortunately, N-body simulations predictions are not supported by observations. Observations of the inner part of density profiles of dwarfs galaxies and LSBs are characterized by a core-like structure (e.g. \\citealt{Flores:94}; \\citealt{Moore:94}; \\citealt{deBlok:02}; \\citealt{deBlok:03}; \\citealt{Gentile:04}, \\citealt{Salucci:07}; \\citealt{deNaray:08}, \\citealt{deNaray:09}), and a similar problem is evidenced when studying the clusters of galaxies inner profile. Density profiles of clusters has been studied through X-ray observations, strong and weak lensing. X-ray temperature measurements give information on cluster structure in the range 500-50 kpc (\\citealt{Brada:08}). At smaller radii, temperature determination is limited by instrumental resolution or substructure (\\citealt{Schmidt:07}, hereafter SA07). X-ray measurements are also limited by ``cooling flows\" presence, and the breaking of assumption of hydrostatic equilibrium (see \\citealt{Arabadjis:04}). The inner slope calculation through the use of X-ray observations brought to discrepant values ( e.g., 0.6 (\\citealt{Ettori:02}), 1.2 (\\citealt{Lewis:03}), 1.9 (\\citealt{Arabadjis:02}).  Another technique used to study the DM distribution in clusters is gravitational lensing. Weak lensing of background galaxies is used to reconstruct the mass distribution in the outer parts of clusters (\\citealt{Mellier:99}). The resolution that can be achieved is able to constrain profiles inside~100 kpc. Strong lensing is used to study the DM distribution in the inner parts of clusters, it has a typical sensitivity to the projected mass distribution inside $\\simeq 100-200$ kpc, with limits at $\\simeq 10-20$ kpc (\\citealt{Gavazzi:05}; \\citealt{Limousin:08}, Alexandrov et al. 2011; Tsvetkova et al. 2009).  Discrepant results have been sometime obtained when using the lensing method. \\citet{Smith:01}, found $\\alpha >1$ at 1\\% of $r_{vir}$ by studying the tangential and radial arcs of A383. A much smaller value $\\alpha$ was obtained by \\citet{Sand:04} and \\citet{Newman:11}, for the same cluster, using lensing and through the aid of stellar kinematics of the central region. \\citet{Tyson:98} found $\\alpha=0.57 \\pm 0.024$ for Cl 0024+1654, while  \\citet{Kneib:03} found that a NFW profile fits the profile in the radius range 0.1-several $r_{vir}$. \\citet{Sand:02} found a cored profile with $\\alpha =0.35$ for MS2137.3-2353, while  \\citet{Gavazzi:03} and \\citet{Gavazzi:05} concluded that the precise value of the slope depends on the mass-to-light ratio of the Brightest Cluster Galaxy (BCG).  In summary, N-body simulations are not always in agreement with the inner slopes of dwarf galaxies, LSBs and clusters of galaxies, and in the case of clusters of galaxies, observations may even disagree for the very same cluster (\\citealt{Gavazzi:05} ; \\citealt{Smith:05}; \\citealt{Zappacosta:06}; SA07; \\citealt{Brada:08}; \\citealt{Limousin:08}; \\citealt{Umetsu:08}).  Several could be the reasons of the discrepancy (e.g., (a) different definition of the slope, which sometimes refers to the DM and sometimes to the total mass; (b) use of observational techniques with different/limited dynamic range in radius; (c) not taking into account the stellar mass of the BCG). In order to obtain more sure constraints on the central part of the density profiles, one should use combined methods (\\citealt{Miralda-Escude:95}; \\citealt{Kneib:03} (weak+strong lensing); \\citealt{Brada:05} (weak+strong lensing); \\citealt{Mahdavi:07} (X-ray+weak lensing)), or better constraints on the central part of the density profiles can be obtained through stellar kinematics of the central galaxy ($\\simeq 1-200$ kpc region).  In a series of papers, \\citet{Sand:02}, \\citet{Sand:04}, \\citet{Sand:08}, studied the clusters MS 2137-23; A383; A963; RXJ1133; MACS 1206; A1201, separating the contribution to the halo coming from the DM from that coming from the baryonic and stellar mass of the BCG. They found a profile flatter than $\\alpha <1$ except for RXJ1133. \\citet{Newman:11} found $\\alpha <1$ (95 per cent CL) for A383. \\citet{Newman:09} presented a detailed analysis of DM and baryonic distributions in A611, finding a slope $\\alpha <0.3$ (68 per cent CL).  Summarizing, at least some clusters of galaxies have inner density-profile slopes shallower that those obtained in N-body simulations, in agreement with what happens, as previously reported, with the density profiles of dwarfs galaxies and LSBs (e.g. \\citealt{Flores:94}; \\citealt{Moore:94}; \\citealt{deBlok:02}; \\citealt{deBlok:03}; \\citealt{Gentile:04}, \\citealt{Salucci:07}; \\citealt{deNaray:08}, \\citealt{deNaray:09}). We want to recall here that, similarly to Sand's result concerning the quoted clusters, the galaxies NGC 2976, NGC 4605, NGC 5949, NGC5963 and NGC 6689, have inner slopes going from very flat to cuspy, and \\citet{deBlok:08}, using a sample from The HI Nearby Galaxy Survey (THINGS) found that the best fit to rotation curves, and then the inner slope $\\alpha$ of their density profile, depends on their mass.  Several papers have shown the fundamental role of baryons in shaping the density profiles of structures. Different processes have been pointed out capable of flattening the inner density profile, transferring energy from stellar baryons to the dark matter, heating it and lowering the central dark matter density (\\citealt{Milosavljevic:01}; \\citealt{El-Zant:01}, \\citealt{El-Zant:01}; \\citealt{Weinberg:02}; \\citealt{Loeb:03}; \\citealt{Gao:07}; \\citealt{Ma:09}; \\citealt{McMillan:05};   \\citealt{Tonini:06}; \\citealt{Mashchenko:06}; \\citealt{Romano-Diaz:08}, \\citealt{Romano-Diaz:09}; \\citealt{delPopolo:09}; \\citealt{Governato:10}; Kulinich et al. 2012).  In the present paper, we use Chandra X-ray data to study, similarly to SA07, the properties of 129 dynamically relaxed galaxy clusters. We selected the sample with redshift range from 0.01 to 1.4. with the aim to recover their total and gas mass profiles and analysis of the measured distribution of $c_{200}$, $M_{200}$ and baryonic mass content. We determine the baryons content of each cluster subtracting the DM from the total mass. Even if the results concerning the density profile of clusters are similar to those of SA07, we point out that fitting the density profile with a generalized NFW model \\begin{equation} \\rho(r)= \\frac{\\rho_0}{(r/r_s)^{\\alpha}(1+r/r_s)^{3-\\alpha}} \\end{equation} the slope $\\alpha$ is correlated with the baryonic mass content of the cluster, in agreement with several studies (\\citealt{Ricotti:03}; \\citealt{Subramanian:00}; \\citealt{Simon:05}; \\citealt{Cen:04}; \\citealt{Ricotti:04}; \\citealt{Williams:04}; \\citealt{Ricotti:07}; \\citealt{delPopolo:09}, \\citealt{delPopolo:12}). In the paper the assumed value of Hubble constant is $H_{0}=73 km s^{-1} Mpc^{-1}$ and $\\Omega_{m}=0.27$. The outline of our work is the following. In Section 2, we describe the X-ray galaxy clusters sample of Chandra observations compiled by us  to recover the baryonic and total mass profiles with techniques presented in Section 3. In Section 4, we present a preliminary discussion of the main results. We investigate the $c_{200}-M_{200}$ and $\\alpha-M_{b}$ relations. We summarize our results and draw the conclusion of the present study in Section 5 and 6. ",
        "conclusions": "We presented the reconstruction of the total mass (dark matter, gas and luminous matter) from the Chandra observations of 129 massive X-ray luminous galaxy clusters in the redshift range 0.01 -- 1.4. We estimated the total ($M_{tot}$) mass within $R_{200}$ in the range 3 -- 293 $\\times 10^{14} M_{\\odot}$.  In order to estimate the fraction of dark matter and gas we have used the conditions of hydrostatic equilibrium and spherical symmetry. Similarly to SA07, we performed two different analysis: the first concerning the total mass of the clusters, and the second decomposing the total mass in its diffuse gas, DM, and BCG components.  Similarly to SA07, the NFW gives a good fit to the total mass and DM distribution. We obtained a best-fitting result for the inner slope of the DM profile in the clusters $\\alpha = 0.94 \\pm 0.13$. We also obtained a mass-concentration relation  $c \\propto M^a/(1+z)^b$, with $a = $ -0.56 $\\pm$ 0.15 and $b = $0.80$ \\pm$ 0.25 (68 per cent confidence limits) in agreement with previous results and simulations.  Finally, we showed that there is a tight correlation among the inner slope $\\alpha$ and the baryonic mass, $M_b$, content."
    },
    "1208/1208.2879_arXiv.txt": {
        "abstract": "\\begin{center} \\begin{minipage}{6.6in} \\qquad \\input{rabstract} \\vskip\\baselineskip \\qquad \\end{minipage} \\end{center} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.2048_arXiv.txt": {
        "abstract": "Key words: signal; noise; gain; quantum efficiency; count; rad noise; dark current; nuisance parameters\\\\ Calibration is nowadays one of the most important processes involved in the extraction of valuable data from measurements. The current availability of an optimum data cube measured from a heterogeneous set of instruments and surveys relies on a systematic and robust approach in the corresponding measurement analysis. In that sense, the inference of configurable instrument parameters can considerably increase the quality of the data obtained.\\\\ Any measurement devoted to scientific purposes contains an element of uncertainty. The level of noise, for example, determines the limit of usability of an image. Therefore, a mathematical model representing the reality of the measured data should also include at least the sources of noise which are the most relevant ones for the context of that measurement.\\\\ This paper proposes a solution based on Bayesian inference for the estimation of the configurable parameters relevant to the signal to noise ratio. The information obtained by the resolution of this problem can be handled in a very useful way if it is considered as part of an adaptive loop for the overall measurement strategy, in such a way that the outcome of this parametric inference leads to an increase in the knowledge of a model comparison problem in the context of the measurement interpretation.\\\\ The context of this problem is the multi-wavelength measurements coming from diverse cosmological surveys and obtained with various telescope instruments. As a first step,, a thorough analysis of the typical noise contributions will be performed based on the current state-of-the-art of modern telescope instruments,  a second step will then consist of identifying configurable parameters relevant to the noise model under consideration, for a generic context of measurement chosen. Then as a third step a Bayesian inference for these parameters estimation will be applied, taking into account a proper identification of the nuisance parameters and the adequate selection of a prior probability. Finally, a corresponding set of conclusions from the results of the implementation of the method proposed here will be derived ",
        "introduction": "\\label{sect:1} As indicated in \\cite{AstronomicalPhotometry:Budding}, astronomical photometry is about the measurement of the brightness of radiating objects in the sky. Many factors, such as those coming from the instrument limitation, from a fixed measurement strategy, or the limitation from the mean through which the measurement is taking place, make this area of the science relatively imprecise. The improvement in the dectectors technology plays a key role in the area of optimizing the resulting astronomical photometric measurements. In this sense, a signal-to-noise ratio capable of being configured as part of an optimization framework of the measurement system seems to be a useful input.\\\\ Charge-coupled devices (CCDs) constitutes the state-of-the-art of detectors in many observational fields. \\cite{CCDHandbook:Howell} enumerates the areas involved in the recent advances of the CCDs systems, which are: \\begin{itemize} \\item Manufacturing standards that provide higher tolerances in the CCD process leading directly to a reduction in their noise output. \\item Increased quantum efficiency, especially in the far red spectral regions. \\item New generation controll electronics with the ability for faster readout, low noise performance, and more complex control functions. \\item New types of scientific grade CCDs with some special properties. \\end{itemize}  Any data, in general, is always limited in accuracy and incomplete, therefore, deductive reasoning does not seem to be the proper way to prove a theory. However, and  as said in \\cite{BayesianPhysicalSciences:Gregory}, statistical inference provides a mean for estimating the model parameters and their uncertainties, which is known as data analysis. It also allows assessing the plausability of one or more competing models.  The use of a Bayesian approach here is also justified in \\cite{BayesianPhysicalSciences:Gregory} where it is stated that for data with a high signal-to-noise ratio for example, a Bayesian analysis can frequently yield many orders of magnitude improvement in model parameter estimation, through the incorporation of relevant prior. This is exactly what we intend through the implementation of what will be described in this paper, and detailed in the following section.   \\section {Description of the problem to be resolved} \\label{sect:2} The problem to be resolved here consists in the implementation of bayesian inference for a set of configurable parameters which affect the signal-to-noise ratio of a measurement. In comparisson with other methodologies, such as ANOVA, which shows serious weakness when outliers are present in the measured data, bayesian parameter inference offers a robust method against outliers. It also allow to improve the results of inference by using the posterior probability density distribution (pdf) of one execution as the prior pdf for another execution in a recursive framework. This will lead to an adaptive measurement strategy which can be addressed as a calibration refinement.  Professional  surveys plan the measurement strategy well in advance, taking into account all the relevant factors impacting on the measurement; this involves the set of specified fix parameters from the detector and also a set of parameters which configure the measurement, such as integration time, diameter of the aperture, etc.  Once a measurement has finished, the data are archived and their analysis and processing begin. The problem proposed here is to establish a link between the results of a measurement under a specific detector configuration and the refinement, by application of parameter bayesian inference,  of the configuration parameters to be applied in a further measurement. The result of this bayesian inference at parameter level is proposed to be injected as additional knowledge for a model selection problem in the context of measurement data analysi (i.e, photometric cross-matching of multiwavelength astronomical sources).  For example, let us imagine that we have performed colour measurement in a multi-wavelength survey  with ten different instruments, each one under a specific configuration. Let us imagine that in the process of model comparison for the cross-matching scenarios, the existence of a source inferred in a bandwith which is not detecting it, is plausible. Then, based on this result, a new configuration for that instrument can be inferred in such a way that allows us to explore the refined plausability indicated by the model comparison from the data obtained in the first measurement loop. Figure \\ref{fig1} shows a  block diagram reflect in general lines the idea of the problem proposed here\\\\  \\begin{figure}[Graficobis] \\label{fig1} \\centering \\includegraphics[width=2.5in]{Graficobis} \\caption{Two levels of Bayesian Inference} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "1208/1208.5497_arXiv.txt": {
        "abstract": "Following two years of complete occultation of both stars by its opaque circumbinary ring, the binary T Tauri star KH~15D has abruptly brightened again during apastron phases, reaching $I = 15$ mag. Here, we show that the brightening is accompanied by a change in spectral class from K6/K7 (the spectral class of star~A) to $\\sim$K1, and a bluing of the system in $V-I$ by about 0.3 mag. A radial velocity measurement confirms that, at apastron, we are now seeing direct light from star~B, which is more luminous and of earlier spectral class than star~A. Evidently, the trailing edge of the occulting screen has just become tangent to one anse of star~B's projected orbit. This confirms a prediction of the precession models, supports the view that the tilted ring is self-gravitating, and ushers in a new era of the system's evolution that should be accompanied by the same kind of dramatic phenomena observed from 1995-2009. It also promotes KH~15D from a single-lined to a double-lined eclipsing binary, greatly enhancing its value for testing pre-main sequence models. The results of our study strengthen the case for truncation of the outer ring at around 4~AU by a sub-stellar object such as an extremely young giant planet. The system is currently at an optimal configuration for detecting the putative planet and we urge expedient follow-up observations. ",
        "introduction": "\\objectname{KH~15D}  is a young binary system composed of similar, but not identical, low-mass pre-main sequence stars in an orbit of eccentricity $\\sim$0.6 with a period of 48.37 days \\citep{johnson04,winn04,hamilton05}. The binary orbit is viewed nearly edge-on and is embedded in an accretion disk from which a well-collimated outflow emerges \\citep{hamilton03,deming04,tokunaga04,mundt10}. The estimated age of the stars is $3\\times 10^{6}$ yr, based on their location in the color-magnitude diagram of NGC 2264, and the total system mass is $\\sim$1.3 M$_\\odot$ \\citep{hamilton01,johnson04}. A thin circumbinary ring of evolved solids has precipitated from the gas disk \\citep{lawler10} within the terrestrial zone (1-4~AU) of this system \\citep{herbst08}. The ring reveals itself as an opaque screen with a razor-sharp edge ($\\lesssim$ 0.1 stellar radii; \\citeauthor{herbst10} 2010) at optical and near-infrared wavelengths. For decades, the leading edge of this structure has been slowly moving across the binary orbit, apparently resulting from precession of the slightly inclined ring \\citep{chiang04,winn04,winn06}. One star's orbit (designated star~B) was completely occulted in 1995 and the other (star~A) in 2009 \\citep{herbst10}.  The tidal influence of a binary on a misaligned disk is known to cause rigid-body precession given that the disk can communicate at length-scales comparable to the inner disk radius and time-scales close to the local precession period \\citep{arty94,pap95,larwood97}. The variable torque applied by a non-coplanar binary will truncate the inner disk radius and induce a gradient of inclinations across concentric annuli in the disk, producing a warp. The specific geometry of the warp, such as whether the gradient increases or decreases with radius, depends upon the mechanism responsible for enforcing an alignment of the nodes of the annuli. \\cite{chiang04} propose two means by which the KH~15D circumbinary ring maintains such an alignment - either by self-gravity of the ring particles or by thermal pressure normal to the ring plane - with a unique warp geometry resulting in each case.  In such warped, precessing disks, if the outer radius is allowed to become arbitrarily large, then the disk precession period will approach infinity. Thus, for nodal precession to occur, the outer disk must also be confined in radial extent (i.e. it forms a small, separate ring), such as by a shepherding body \\citep{ goldreich82, chiang04} or a sufficiently rapid decrease in surface density \\citep{larwood97}. The misalignment of the binary orbital plane with the KH~15D ring is believed to induce precession in its ring and, further, the system exhibits distinctive eclipses due to this non-coplanar arrangement.  Although identified as a variable star in the late 1960s \\citep{badalian70}, interest in KH~15D began with the discovery in 1995 \\citep{kearns98} that the star essentially winked on and off by several magnitudes on a period of 48.37 days \\citep{hamilton05}, now known to be the orbital period of the binary \\citep{johnson04}. These extreme brightness excursions started then because precession of the ring had moved one of its razor-sharp edges fully across the orbit of one member of the binary (star~B), leaving the other (star~A) to periodically rise and set with respect to that horizon as seen from Earth. From 1995 to 2009, the winking continued, with the system being increasingly more ``off'' as star~A spent more and more of each orbital cycle below the horizon \\citep{hamilton05,herbst10}. In 2010 and 2011, the system was fainter at all times because neither star appeared above either edge.  Both stars are believed to be less massive than the Sun, with star~A having a spectral type of K6/K7 \\citep{hamilton01,agol04}. The prediction that star~B is slightly more massive and of earlier type has heretofore only been inferred from the fact that it is a brighter star, as evidenced by archival photometry \\citep{winn03,johnsonwinn04,johnson05} and by its one appearance in the modern photometric record (see Fig.~\\ref{longterm}) prior to 2012. ",
        "conclusions": "\\subsection{Confirmation of the Precession Model}  The precessing circumbinary-ring model has provided a successful framework for interpreting observations of KH~15D \\citep{winn04,winn06,chiang04,silvia08}. A prediction is that the eclipse behavior exhibited by star~A from the early 1990s through 2009 should repeat itself with star~B as the visible star. The developments reported here confirm that prediction and provide strong support for this scenario. Fig.~\\ref{screenmotion} is a schematic drawing of the system, as we believe it appears in 2012. For the first time, we are able to locate the trailing edge of the ring which provides important new information to improve the next generation of models.  \\cite{winn04,winn06} employed archival records to address the secular evolution of KH~15D by projecting the binary orbit onto the plane of the sky and treating the precession of the ring as a sharp, semi-infinite edge (that is, the location of leading edge is finite and the trailing edge extends to infinity) that marches across the projected binary orbit over time. The authors predicted the eventual re-emergence of star~B as a natural result of the advancing edge. As the model had no way to differentiate between viewing angles with respect to the occulting feature nor its thickness, it provided no definite prediction regarding whether both binary orbits would ever be occulted at the same time or, if so, when star~B  might reappear.  Having the location of only one edge of the screen specified also complicated the modeling of the scattered light. The discovery of the location of the trailing edge should lead to a much improved version of this model which will ultimately quantify most of the important physical characteristics of the system.  \\cite{chiang04} presented a first-order dynamical theory of the KH~15D ring noting that its rigid precession can be enforced either by thermal pressure gradients or by ring self-gravity. They favored self-gravity and their case is strengthened by the observations contained herein. The future light curve of KH 15D predicted by \\cite{chiang04} for the thermal pressure case (Model~1) differed substantially from that for self-gravity (Model~2) because the direction of warping is different. In Model~1, the disk is closer to the plane at large distances and more inclined closer in and the predicted duration of full occultation is quite long (many decades). In Model~2, the disk is more inclined at larger distances and the duration of full occultation of both orbits is short or non-existent. We now know that the observed light curve had a short duration ($\\sim$2 years) of full occultation, qualitatively in much better agreement with the self-gravity model, although quantitatively in agreement with neither.  While recent data will enable properties such as the ring geometry, orientation and scattering properties to be refined, the mounting support for rigid nodal precession underscores that the occulting ring of solids must be confined at its outer borders. Therefore, the inferred radial dimensions of the ring should not change drastically from past estimates.  \\subsection{Remaining Questions and Future Observations}  The question remains as to what has truncated the ring at around 4~AU. \\citet{chiang04} suggest a planet as the most likely cause. Being at or near the ice line in the system it is possibly a giant planet in formation. Color excursions to values redder than star~A have long been known in KH~15D during minimum light and may find their interpretation in the light of this putative third body.  Such an extremely young giant planet, if sufficiently massive, might be self-luminous enough to be detectable in the near or mid-infrared during phases when both stars are occulted. It should be sought immediately, as  the minimum brightness of the system is already beginning to gradually increase due to the advancement of the trailing edge (see Fig.~\\ref{longterm}).  Fig.~\\ref{last3lightcurves} shows that the abrupt change in slope that occurs during star rise and star set continues, suggesting that the trailing edge of the ring will be as sharp as the leading edge: namely, much smaller than a stellar radius \\citep{herbst10}. Resultantly, the eclipses that will ensue will allow us to determine the radius of star~B with the same precision that can be achieved in a typical eclipsing binary. We can now obtain the RV curve of star~B, and so will constrain the masses, radii, luminosities and effective temperatures of the stars. KH~15D will soon be the first known double-lined eclipsing binary young enough to still be embedded in an accretion disk, making it of paramount importance for testing models of pre-main sequence evolution.  As Fig.~\\ref{screenmotion} shows, the physical extent of the obscuring portion of the ring is slightly wider than the projected size of the binary orbits. What sets this width remains unclear to us. Multiple possibilities exist but it is well beyond the purview of this observational paper to assess them here. The KH~15D spectral features of non-stellar origin will be addressed in subsequent work, and additional spectra should be obtained to verify these features and probe more sight lines in the ring.  Over the coming decade we will have the opportunity to use the trailing edge for natural coronography to study the magnetosphere and jet-launching zone of star~B in the same detail as was done for star~A."
    },
    "1208/1208.3614_arXiv.txt": {
        "abstract": "Creations of light anti-nuclei (anti-deuterium, anti-tritium, anti-He3 and anti-He4) are observed by collaborations at the LHC and RHIC accelerators. Some cosmic ray experiments are aimed to find the anti-nuclei in cosmic rays. To support the experimental studies of the anti-nuclei a Monte Carlo simulation of anti-nuclei interactions with matter is implemented in the {\\sc Geant4} toolkit. The implementation combines practically all known theoretical approaches to the problem of antinucleon-nucleon interactions. ",
        "introduction": "\\label{intro} One of the most exciting puzzles in cosmology is connected with the question of the existence of anti-matter in the Universe. A number of dedicated cosmic ray experiments aim to search for anti-nuclei \\cite{PAMELA,BESS,AMS,CAPRICE}.  Also, anti-nuclei have been observed in nucleus-nucleus and proton-proton collisions  by experiments at the RHIC \\cite{STAR,PHENIX,PHOBOS} and LHC accelerators \\cite{ALICE}. The STAR collaboration at RHIC reported in March 2011 that the anti-He4 nuclei were identified in high energy nucleus-nucleus collisions (see \\cite{StarAHe4}). The ALICE collaboration at LHC confirmed the STAR results in May 2011.  An experimental study of anti-nuclei requires a knowledge of anti-nucleus interaction cross sections with matter. The cross sections are needed to estimate various experimental corrections, especially those due to particle losses which reduce the detected rate. In practice, various phenomenological approaches are applied in order to estimate the antinucleus-nucleus cross sections. Thus, a first task is a creation of reliable estimations of the cross sections. Here we use the Glauber approach.  It is obvious that an annihilation can take place at an interaction of an anti-nucleus with a nucleus. A lot of mesons can be produced in this way. Thus, we have to simulate the meson production. We do this in the framework of the quark-gluon string model.  Low energy mesons can have secondary interactions in nuclear residues. We take them into account using the binary cascade model of the {\\sc Geant4} toolkit \\cite{GEANT4}. ",
        "conclusions": ""
    },
    "1208/1208.4034_arXiv.txt": {
        "abstract": "We present a 3-dimensional analysis of the supernova remnant Cassiopeia~A using high resolution spectra from the Spitzer Space Telescope.  We observe supernova ejecta both immediately before and during the shock-ejecta interaction.  We determine that the reverse shock of the remnant is spherical to within 7\\%, although the center of this sphere is offset from the geometric center of the remnant by 810 km s$^{-1}$.  We determine that the velocity width of the nucleosynthetic layers is $\\sim$1000 km s$^{-1}$ over ~4000 square arcsecond regions, although the velocity width of a layer along any individual line of sight is $<$250~km~s$^{-1}$.  Si and O, which come from different nucleosynthetic layers in the progenitor star, are observed to be coincident in velocity space in some directions, but segregated by up to $\\sim$500 km s$^{-1}$ in other directions.  We compare these observations of the nucleosynthetic layers to predictions from supernova explosion models in an attempt to constrain such models.  Finally, we observe small-scale, corrugated velocity structures that are likely caused during the supernova explosion itself, rather than hundreds of years later by dynamical instabilities at the remnant's reverse shock. ",
        "introduction": "The supernova remnant Cassiopeia A (Cas~A) is a unique astrophysical laboratory due to its young age \\citep[$\\sim$340 years -][]{thor01, fesen06} and small distance \\citep[only 3.4 kpc -][]{reed95}.  The remnant is just entering its Sedov-Taylor phase, so emission from both forward and reverse shocks can be detected \\citep{hughes00}.  Emission at most wavelengths, including most of the infrared, is dominated by a $\\sim$120$\\arcsec$ radius ``Bright Ring''.  The Bright Ring is formed when supernova ejecta encounter Cas~A's reverse shock and are shocked, heated, and collisionally ionized.  It consists of undiluted ejecta rich in O, Si, S, Ne, Ar, Ca, and Fe \\citep{chev78,douv99,hughes00,will03,hwl03,lhw03,morse04,ennis06}.  Studies of optical light echoes from the the explosion near peak light have led to the observation of weak hydrogen lines, indicating a supernova Type IIb origin for Cas~A \\citep{krause08}.  Cas~A's progenitor was therefore a red supergiant that had lost most, but not all, of its hydrogen envelope.  X-ray studies indicate a total ejecta mass of $\\sim$2M$_{\\odot}$ \\citep{will03}. If one adds to this the mass of the central compact object \\citep{chak01}, Cas~A's progenitor had a total mass of at least 4M$_{\\odot}$ immediately before the supernova explosion.  The estimated oxygen mass indicates a main sequence mass of $\\sim$15-25M$_{\\odot}$ \\citep{young06, vink96}.  Although Cas~A's appearance is dominated by recently shocked ejecta, it also contains emission that is not the result of collisional ionization at the reverse shock, but photoionization by UV and X-ray emission from the shocked ejecta \\citep{hs84, hf98, smith09}.  This material is seen toward the central region of the remnant at low radio frequencies \\citep{kas95} and in the infrared \\citep{rho08, smith09, delaney10, isens10}, and was demonstrated to be at lower densities and ionization state than recently shocked material on the Bright Ring through a combination of Doppler analysis and line ratio measurements \\citep{smith09}.  These ejecta are often referred to as ``unshocked ejecta'' since they have yet to encounter the remnant's reverse shock.  That is not an accurate label, since Cas~A's forward shock and a reverse shock interacted with the ejecta during the supernova explosion itself.   \\subsection{Previous 3D Maps}  3D Maps of Cas~A have been made in the optical, infrared, and X-ray.  Doppler reconstructions in the optical used S and O emission lines \\citep{law95,reed95} and showed that ejecta on the Bright Ring lie on a roughly spherical shell but are not uniformly distributed on that shell - most of the ejecta lie nearly in the plane of the sky.  They also observe that the center of the sphere is offset from the geometrical center of the spherical shell by $\\sim$0.36pc along our line of sight.  This indicates that the ejecta are not traveling at the same velocity in all directions, which is consistent with previous results which indicated an asymmetric expansion for the ejecta \\citep[e.g.][]{braun87, will02}.  These 3D reconstructions give us a selective snapshot of ejecta because only material that has recently encountered the remnant's reverse shock will emit strongly in the optical.  \\cite{delaney10} created a 3D infrared and X-ray map of Cas~A from a \\emph{Spitzer Space Telescope} spectral cube\\footnote{Movies showing this 3D structure are available at http://chandra.harvard.edu/photo/2009/casa2/animations.html}.  \\cite{isens10} used a similar IR data set, but at higher spectral resolution, to make a 3D map of ejecta in the center of the remnant.  The advantage of these IR maps lies in the fact that much of the ejecta in the IR will be detectable both before and after they interact with the reverse shock.  Both studies found a similar distribution of ejecta to that seen in the optical where the center of expansion is offset from the geometrical center of the remnant both in projection and along the line of sight.  These works were able to study the relationship of several nucleosynthetic layers and are discussed in the next section.  \\subsection{Separation of Nucleosynthetic Layers}  Si and O emission are observed to be co-located in most regions \\citep[e.g.][]{ennis06} in both the X-ray and infrared.  This indicates that the two layers have very similar velocities (less than 80 km~s$^{-1}$ difference).  However, evidence of layer differentiation is found in some directions in the X-ray \\citep[e.g.][]{hughes00}, the optical \\citep{fesen06}, and the IR \\citep[e.g.][]{delaney10,isens10}, which was likely caused by the different layers of the star being ejected at different velocities in those directions, thus encountering the remnant's reverse shock at different times.     It should be emphasized that we can only observe mixing or separation in \\emph{velocity space}.  We can easily detect any velocity gradients in the supernova explosion since we can detect Doppler velocities of $<$100~km~s$^{-1}$ in the IR, while typical observed velocities and velocities predicted by models are an order of magnitude larger \\citep[e.g.][]{hammer10}.  However, we cannot detect any initial spatial separation of the nucleosynthetic layers - simulations predict that the relevant nucleosynthetic layers will be $<$ $10^{11}$ cm thick prior to the explosion \\citep[e.g.][]{jog09}, but the typical ejecta clump size of $<$1$\\arcsec$ corresponds to $\\sim$ $10^{16}$ cm at Cas~A's distance.  Therefore, we cannot differentiate between a situation where two nucleosynthetic layers were separated during the supernova explosion but ejected at the same velocity, and one where the two layers were completely mixed during the explosion and ejected at the same velocity.  But, if we observe two layers that are currently separated in velocity space, we know that they were separated during the supernova explosion itself because, to the best of our knowledge, there is no mechanism that will impart different velocities to spatially overlapping elements.     \\subsection{Geometrical Asymmetries}  Supernova explosion models predict substantial asymmetries due to effects such as rotation as well as instabilities \\citep[e.g.][]{blondin03, burrows07, hammer10}.  Observations of both supernovae and supernova remnants have confirmed this picture.  Spectropolarimetric observations of unresolved supernovae have shown that all observed core collapse supernovae contain intrinsic polarization, indicating that there is a departure from spherical symmetry \\citep{whe05}.  Although an axis-symmetric geometry, probably induced by jets, can be used to explain some features in some core collapse supernovae, significant departures from axial symmetry are needed to explain most observations \\citep{wang08}.  IR 3D maps of supernova ejecta in Cas~A have found major asymmetries, both on global scales \\citep{delaney10} and for smaller subsets of ejecta in the the entire supernova remnant \\citep{isens10}.   These asymmetries are not immediately apparent in the visual appearance alone because the highly spherical reverse shock creates a large selection effect in that we can only observe ejecta near the shock at most wavelengths.  In this paper, we present an analysis of high spectral resolution Spitzer mappings of the ejecta on the Bright Ring of Cas~A.  This data set is an extension of that used by \\cite{isens10}, and it contains regions with both recently shocked and interior ejecta.  In $\\S$2 we present the observations and discuss the methods used in our analysis.  We describe those results in $\\S$3 and discuss the physical implications in $\\S$4. ",
        "conclusions": "We create a 3D model of shocked ejecta of Cas~A in select regions at unprecedented spectral resolution using IR ionic lines.  We confirm previous studies that indicate that the remnant is offset by $\\sim$800 km~s$^{-1}$ along our line of sight.  We find evidence for velocity separation between the O and Si layers along some, but not all, lines of sight.  We measure the velocity width of these layers roughly 250 km~s$^{-1}$ thick for a single line of sight, although the ejecta are often in bands that are $\\sim$1000 km~s$^{-1}$ thick averaged over several nearby lines of sight due to corrugation.  We find evidence for corrugation in some regions of the remnant, and speculate that the corrugation was caused during the explosion itself rather than hundreds of years later.  We use our observations of Si and O velocities to begin constraining models of supernova explosions, and to motivate future models to explore velocity profiles as a function of azimuth.  We look forward to potential similar data sets from instruments such as the Herschel Space Observatory and the Stratospheric Observatory for Infrared Astronomy (SOFIA).  Both these observatories will have the ability to create spectral cubes of Cas~A at much higher spectral resolution.  The current instruments on both observatories do not have the necessary instantaneous bandwidth to observe the 10,000 km~s$^{-1}$ velocity range of ejecta in Cas~A, but future spectrographs will have scanning modes that will allow observations of ejecta at many different velocities."
    },
    "1208/1208.5338_arXiv.txt": {
        "abstract": "We investigate the temporal and spectral correlations between flux and anisotropy fluctuations of TeV-band cosmic rays in the light of recent data taken with IceCube. We find that for a conventional distribution of cosmic-ray sources the dipole anisotropy is higher than observed, even if source discreteness is taken into account. Moreover, even for a shallow distribution of galactic cosmic-ray sources and a reacceleration model, fluctuations arising from source discreteness provide a probability only of the order of 10\\% that the cosmic-ray anisotropy limits of the recent IceCube analysis are met. This probability estimate is nearly independent of the exact choice of source rate, but generous for a large halo size. The location of the intensity maximum far from the Galactic Center is naturally reproduced. ",
        "introduction": "The anisotropy in the arrival directions of cosmic rays in the TeV band has received renewed attention in recent years. A variety of experiments \\citep{2006Sci...314..439A,2007PhRvD..75f2003G,2009ApJ...692L.130A, 2009ApJ...698.2121A,2010cosp...38.2707Z,2010ApJ...718L.194A,2012ApJ...746...33A, 2012arXiv1202.3379D} have reported results on the amplitude of the first harmonic in the siderial anisotropy, which is typically found to be slightly less than $0.1\\%$. All experiments see only part of the sky, and it is not easy to reconstruct an allsky dipole anisotropy, largely because the energy dependence of the acceptance is different among the various detectors. In any case, the true dipole anisotropy can be slightly higher than the sidereal first harmonic, i.e. $\\lesssim 0.1\\%$. In addition, significant small-scale anisotropy was observed that has not met an accepted explanation to date, although it is tempting to attribute it to the local structure of the turbulent magnetic field in the Galaxy \\citep[e.g.][]{2011arXiv1111.2536G,2012PhRvL.108z1101G}.  Here we investigate whether or not the dipole anisotropy can be reproduced with models of galactic cosmic-ray propagation that are tuned to fit available data on secondary-to-primary ratios and the survival fraction of unstable isotopes. For that purpose we use a Green-function method that solves a time-dependent diffusion equation, similar to \\citet{2012JCAP...01..011B}. Diffusion equations of this type are derived by averaging the original Fokker-Planck transport equation for the isotropic part of the cosmic-ray distribution function. By construction, anisotropy in this treatment arises only from the diffusive (and possibly drift or convective) flux and is always dipolar in nature. Modeling the small-scale anisotropy, on the other hand, would be described with a transport equation derived from the original Fokker-Planck equation for the anisotropic part of the cosmic-ray distribution and is not considered here.  Our time-dependent treatment permits accounting for a discrete nature of cosmic-ray source. A similar technique was used to model variations in the spectrum of cosmic-ray electrons \\citep{pe98,2003A&A...409..581P, 2011JCAP...02..031M} and cosmic-ray ions \\citep{2005ApJ...619..314B,2006AdSpR..37.1909P, 2012JCAP...01..010B,2012JCAP...01..011B}. We test various cosmic-ray source distributions in galactocentric radius, using the appropriate transport parameters, and compute the expected temporal evolution of the anisotropy. This approach permits a realistic assessment of the likelihood to find a certain anisotropy for a specific choice of parameters. The occasionally discussed mean variance of the anisotropy is a considerably less useful quantity, because the distribution of anisotropy amplitudes is highly skewed.  We are particularly interested in the temporal and spectral correlations between flux and anisotropy fluctuations. Our approach is ideally suited to recover these correlations and investigate their effect, when we compute flux and anisotropy over a finite spectral acceptance and compare with observational results. ",
        "conclusions": "  \\begin{itemize} \\item Consistent with earlier findings, a cosmic-ray source distribution in the Galaxy that fits to the deduced populations of pulsars or SNR, combined with propagation parameters that reproduce the observed secondary-to-primary ratios, will lead to a dipole anisotropy that is higher than observed, even if source discreteness is taken into account. \\item Also in agreement with earlier studies, a shallow energy dependence of diffusion is necessary, thus requiring a reacceleration model to fit the GeV-band Boron-to-Carbon ratio. \\item For a shallow source distribution, fluctuations arising from source discreteness provide a probability of the order of 10\\% that the cosmic-ray anisotropy limits in the high-energy and low-energy bands of the recent IceCube analysis can be simultaneously met, approximately independent of the exact choice of source rate, but less likely so for a large halo size. \\item The location of the intensity maximum far from the Galactic Center and at different locations for the 2 energy bands are naturally reproduced when source discreteness is taken into account. \\end{itemize}  It is unclear what type of cosmic-ray source in the Galaxy would have a distribution in the Galaxy as shallow as required in our study. The question arises whether or not modifications to the description of diffusive transport can render a steep source distribution viable. The critical constraint appears to be imposed by very low dipole anisotropy in the high-energy band of IceCube data, i.e. at approximately 100~TeV energy per nucleon. Is it possible that the extrapolation to this energy of the cosmic-ray propagation parameters does not hold?  Recent data suggest that cosmic-ray spectra \\citep{2010ApJ...714L..89A,2011Sci...332...69A}, and the Boron-to-Carbon ratio \\citep{2012ApJ...752...69O}, flatten above a few hundred GeV/nuc. At this time we do not have a prevalent interpretation of both the cosmic-ray spectral hardening and the flat B/C ratio above a few hundred GeV/nuc. The spectral hardening of cosmic-ray primaries could simply be a fluctuation arising from source discreteness, but in that case the Boron-to-Carbon ratio should fall off more steeply, because the flux of cosmic-ray secondaries fluctuates very little \\citep{2005ApJ...619..314B}. Spectral hardening can also reflect variations in the production spectra of different sources \\citep{1972ApJ...174..253B,2001A&A...377.1056B,2011PhRvD..84d3002Y}, in which case the Boron-to-Carbon ratio would be unaffected. If the upturn in the Boron-to-Carbon ratio is, on the other hand, due to the fragmentation of primaries and the re-acceleration of secondaries inside cosmic-ray sources \\citep{2009PhRvL.103h1104M,2012A&A...544A..16T}, then the energy dependence of the diffusion coefficient cannot be argued to be very shallow. The only possibility that would imply a very low anisotropy is a change of the energy dependence of the diffusion coefficient at a few hundred GeV/nuc, e.g. from $\\delta=0.25$ to $\\delta=0.15$ \\citep{2012ApJ...752...68V,2012A&A...544A..16T}. Speculative though that may be, the extrapolation using $\\delta=0.15$ to $10^{17}$~eV may render viable the notion that a significant fraction of the sub-ankle cosmic-ray flux is galactic in origin \\citep[e.g.][]{2011ApJ...742..114P}.  Recently, \\citet{2012PhRvL.108u1102E} suggested a common solution to the anisotropy and cosmic-ray gradient problems, that involves a strong toroidal magnetic field in the Galaxy. Diffusion perpendicular to the azimuthal magnetic field (with coefficient $D_\\perp\\propto E^{0.6}$, as opposed to $D_\\parallel\\propto E^{0.3}$) would be the dominant transport mechanism, assuming field-line wandering and large-scale velocity turbulence can be ignored. The authors further posit that $D_\\perp (r_{\\rm GC})\\propto 1/D_\\parallel(r_{\\rm GC})$ and $D_\\perp\\propto Q(r_{\\rm GC})^\\tau$, where $Q(r_{\\rm GC})$ is the density of cosmic-ray sources and $\\tau \\simeq 0.85$. At least up to 1 TeV particle energy, the predicted radial anisotropy appears to be commensurate with measured values.  We posit that the scenario would do little, if anything, to resolve the 3-D anisotropy problem for discrete sources (the authors only discuss anisotropy for continuous sources in a 2-D scenario). In the direction of fast diffusion,i.e. along the large-scale toroidal magnetic field, we see particles of the same age from sources at larger distance, thus increasing their contribution to the anisotropy. The volume sampled, and thus the number of sources contributing to the local cosmic-ray flux, will also increase, thus reducing the likelihood of a fluctuation toward low anisotropy, but not efficiently enough to overcome anisotropy arising from source discreteness. In the scenario of \\citet{2012PhRvL.108u1102E} we therefore expect significantly more anisotropy in direction of the ordered magnetic field than in radial direction."
    },
    "1208/1208.0171_arXiv.txt": {
        "abstract": "{} {The intensity profiles of the C$_2$ Swan bands in cool DQ white dwarfs cannot be adequately fitted with models that otherwise succesfully reproduce spectral features of the molecule CH in these stars. Initial modelling showed that a two{-}component atmosphere in the style of a spot { might be able to} solve the problem.} {We photometrically observed the two cool DQ white dwarfs GJ1117 and EGGR78 to search for variability caused by stellar spots.} {We have not found any such variability, but we estimate the effects of hypothetical spots on lightcurves. We also estimate detection probabilities for spots in different configurations. Alternative explanations of the problem are needed and briefly discussed. } {} ",
        "introduction": "Magnetism in { white dwarfs (WDs)} is usually studied through Zeeman splitting of atomic spectral lines. This method has proven to be very useful, and with the sensitivity added by spectropolarimetry it is very effective in detecting magnetic fields. {Spectra of} cool DQ WDs do not show any atomic lines, { therefore} an alternative method is required. Fortunately, the carbon molecules present in { most of these stars} are sensitive to magnetic fields and they can be studied using spectropolarimetry.   { During the past several years w}e have used circular spectropolarimetry to observe known DQ WDs to look for polarization signals from C$_2$ and CH molecules that can be found in their atmospheres. We have used two different telescopes, the Nordic Optical Telescope (NOT) { on} the Canary Islands and the ESO VLT (Cerro Paranal, Chile). So far, we have observed 12 objects and found one of them (GJ841B) to be magnetic \\citep{vor10}. The rest of the stars do not show any polarization signal at the { noise} level of our observations, i.e. 0.5 \\% { in Stokes V/I} for the NOT observations and 0.2 \\% for the VLT observations.  In \\citet{vor10} we showed that the model presented in \\citet{ber05, ber07} works very well for the CH absorption bands in GJ841B and were also successful in modelling the CH and C$_2$ blend at 430 nm visible in both G99-37 and GJ841B. But while we modelled the intensity profiles of the purely C$_2$ Swan bands, we were unable to achieve good fits. Problems with fitting Swan bands in DQ WDs have been reported before (for example, \\citealt{duf05} have { shown} that the $\\Delta v=+2$ bands are consistently too strong in their model fits). Our model shows the same problem (as seen { in} Fig. \\ref{fig:gj893}), but in addition to that the $\\Delta v=0$ band appears too weak. Hints of a too weak $\\Delta v=0$ band can also be seen in \\citet{duf05}. These problems persisted for all non-magnetic DQ WDs in our survey, as long as we used a single temperature model. \\begin{figure} \\resizebox{\\columnwidth}{!}{\\includegraphics{GJ893_spotless.eps}} \\caption{Model fit from \\citet{ber05} applied to GJ893, one of the non-magnetic objects in our survey. Transition designations are given above the spectrum. The discrepancy between the depths of different Swan bands is obvious. The temperature of the model is 6000K.} \\label{fig:gj893} \\end{figure}  We started searching for a solution to this problem by combining two models in { the} manner of a photosphere with a spot to { obtain}  reasonable fits to the spectra. This led to somewhat strange results. The best{ -}fitting models usually had a small spot with $T=8000-10000$ K and a photosphere of $T=2000$ K. This would suggest a cool disk or an envelope of gas or dust around the WD with a hole in it through which the atmosphere would be seen. Although dust disks have been found around hot WDs and the central objects of planetary nebulae \\citep{bil11}, this does not sound like a very probable scenario for our cool objects. Instead, { a cool spot has} been found before on a weakly (70 kG) magnetic WD \\citep{bri05}.  To exclude spots or envelopes with holes as the reason for the observed properties in these WDs, we started a photometric observing { programme} with a small remotely controlled telescope on { the} Canary Islands. Since WDs are rotating just like any other star, the movement of the spot on the visible hemisphere would cause some photometric variability. We report the results of our photometric observations here and discuss what we can deduce from them. At the end of the paper we discuss other possible reasons for troubles with modelling. ",
        "conclusions": "We monitored two cool DQ WDs to search for photometric variability as a { signature} of a stellar spot. We did not find { any} such variability. The accuracy and timing of our observations do not allow us to detect very long rotational periods (of { about} a year or longer), very small spots, or some spot configurations. Previous studies of WDs have come to the conclusion that the rotation periods usually { last} from hours to decades \\citep[][and references therein]{cha09}. Some magnetic WDs have indeed been found to be slow rotators with periods of about 100 years \\citep{schnor91, ber99}. Since we did not detect any intra-night variability, the shorter end of the period distribution can be ruled out. Based on these arguments, we can assume that GJ1117 and EGGR78 still might have very long rotational periods.   Although we { cannot} rule out spots on these WDs completely, we consider the values given by our model for the spot sizes and temperatures { to be} very { unrealistic}. But if the spot model is not the solution to the discrepancy in intensity in the carbon molecular bands, { what} better hypotheses { could be found}? We think the answer may lie in the  oscillator strengths of carbon molecules { that are}  embedded in our model. These parameters determine { the strengths of} the individual absorption lines within the molecular bands. By modifying these values slightly, we hope to achieve consistent fits to all molecular bands without invoking a spot model. The current values { were} determined in laboratory conditions, but still contain significant error limits.  If modifying the oscillator strengths will not work, we have to go deeper into the physics of the problem and investigate { in which way} the structure, and therefore, the properties of a C$_2$ molecule is different { in} a white dwarf from conditions in { the} laboratory where the molecule { was} studied. Or maybe our treatment of the WD atmosphere is { incorrect} and we have to consider its properties more thoroughly."
    },
    "1208/1208.5879_arXiv.txt": {
        "abstract": "{Little is known {about} the stellar environment and the genealogy of our solar system. Short-lived radionuclides (SLRs, mean lifetime $\\tau$ {shorter} than 100 Myr) that were present in the solar protoplanetary disk 4.56 Gyr ago could potentially provide insight into that key aspect of our history, were their origin understood.} {Previous models failed {to provide} a reasonable {explanation of} the abundance of two key SLRs, $^{26}$Al ($\\tau_{26}$ = 1.1 Myr) and $^{60}$Fe ($\\tau_{60}$ = 3.7 Myr), at the birth of the solar system by requiring unlikely astrophysical conditions.  Our aim is to propose a coherent and generic solution based on the most recent understanding of star-forming mechanisms.} {% Iron-60 in the nascent solar system is shown to have been produced by a diversity of supernovae belonging to a first generation of stars in a {giant molecular cloud}.  Aluminum-26 is delivered into a dense collected shell by a single massive star wind belonging to a second star generation. The Sun formed in the collected shell as part of a third stellar generation. Aluminum-26 yields used in our calculation are based on new rotating stellar models in which $^{26}$Al is present in stellar winds during the star main sequence rather than during the Wolf-Rayet phase {alone}.   % Our scenario eventually constrains the time sequence of the formation of the two stellar generations {that} just preceded the solar system formation, {along with} the number of stars born in these two generations.} {We propose a generic explanation for the past presence of SLRs in the nascent solar system, based on a collect-injection-and-collapse mechanism, occurring on a diversity of spatial/temporal scales. In that model, the presence of SLRs with a diversity of mean lifetimes in the solar protoplanetary disk is simply the fossilized record of sequential star formation within a hierarchical interstellar medium (ISM). We identify the genealogy of our solar system's three star generations {earlier. } In particular, we show that our Sun was born together with a few hundred stars in a dense collected shell situated at a distance of 5-10 pc {from} a parent massive star having a mass {greater} than about 30 solar masses and belonging to a  cluster containing $\\sim$ 1200 stars.} ",
        "introduction": "Short-lived radionuclides (SLRs) are radioactive elements with mean lifetimes {under} 100 Myr {that} were incorporated {into} meteorites' primitive components such as calcium- {and} aluminum-rich inclusions (CAIs) or chondrules during the earliest evolution phases of our solar system. Understanding their origin has long been a major goal of cosmochemistry \\citep{Russell2001}, as it is essential for constraining the stellar environment of the Sun at its birth \\citep{Meyer2000}, as well as for establishing a chronology of the solar system{'s} first million years \\citep{Keegan04}.   {The} SLRs with the longest mean lifetimes ($\\gtrsim$ 5 Myr, such as $^{129}$I [$\\tau_{129}$ = 23.5 Myr] or $^{244}$Pu [$\\tau_{244}$ = 115 Myr]) have abundances compatible with that of the expected Galactic background {owing} to continuous star formation on kpc spatial scales and tens of Myr timescales \\citep{Huss2009}. {Those} SLRs with shorter mean lifetimes ($\\lesssim$ 5 Myr) appear to be in excess relative to that same background abundance. Among these, $^{10}$Be, $^{36}$Cl{,} and $^{41}$Ca can be made by solar energetic particle irradiation of the protoplanetary disk \\citep{Gounelle2006,Duprat2007, Jacobsen2011}. For $^{10}$Be {alone}, {both} an interstellar origin \\citep{Desch2004} {and} a solar-wind implantation model have also been evoked \\citep{Bricker2010}. The origin of two key SLRs, $^{26}$Al ($\\tau_{26}$ = 1.1 Myr) and $^{60}$Fe ($\\tau_{60}$ = 3.7 Myr), remains elusive.   After decades of measurements within CAIs (the first solids to have formed in our protoplanetary disk), the solar system{'s} initial $^{26}$Al/$^{27}$Al ratio is well established at 5.3 $\\times$ 10$^{-5}$ \\citep{MacPherson1995,Jacobsen2008}, though its homogeneity is still subject to debate \\citep{Villeneuve2009,Liu2012,Gounelle2005}. The situation is a bit more complicated for $^{60}$Fe {because} its record in CAIs is hampered by secondary processes {and} by nickel nucleosynthetic anomalies  \\citep{Birck1988,Quitte2007}, whose origin is far from being understood   \\citep{Steele2011}. Nickel-60 excesses attributed to the decay of  $^{60}$Fe were found in chondrules  \\citep{Tachibana2003,Telus2011}, which are believed to have formed $\\sim$1 Myr  after CAIs \\citep{Villeneuve2009}. Based on these data, the presently accepted upper limit of $^{60}$Fe/$^{56}$Fe is $\\sim$ 3 $\\times$ 10$^{-7}$\\citep{Dauphas2008, Gounelle2010, Telus2011}.   To calculate the $^{60}$Fe and $^{26}$Al concentrations at the onset of solar system formation, in addition to the measured ratios presented above, we rely on the {\\it protosolar} abundances given by Lodders (2003). With $^{56}$Fe/$^{1}$H = 3.2 $\\times$ 10$^{-5}$, $^{27}$Al/$^{1}$H = 3.5 $\\times$ 10$^{-6}$ and a hydrogen mass fraction of 0.71 \\citep{Lodders2003}, we obtain the following concentrations in the nascent solar system for $^{26}$Al and $^{60}$Fe: C$_{\\sun}$ [$^{26}$Al] = 3.3 $\\times$ 10$^{-9}$ M$_{\\sun}$/M$_{\\sun}$ = 3.3 ppb (parts per billion) and C$_{\\sun}$ [$^{60}$Fe] = 4.0 $\\times$ 10$^{-10}$ M$_{\\sun}$/M$_{\\sun}$ = 0.4 ppb. The initial $^{26}$Al/$^{60}$Fe mass ratio was thus equal to 8.2. Such elevated concentrations of $^{26}$Al and $^{60}$Fe {need} to be explained, and are the subject or the present paper.  Though Asymptotic Giant Branch (AGB) stars have been proposed as a possible source of SLRs \\citep{Wasserburg2006,Trigo2009,Lugaro2012}, massive stars (M $\\ge$ 8 M$_{\\sun}$) are the best candidates to account for {the presence of} $^{26}$Al and $^{60}$Fe in the nascent solar system. This is because massive stars at all stages of their evolution are present in star-forming regions, {unlike} AGB stars \\citep{Kastner1994}.  The most massive stars (M $\\gtrsim$ 25 M$_{\\sun}$)  burn hydrogen for million{s of} years on the main sequence (MS) before they enter the short-lived Wolf-Rayet (WR) phase that precedes the supernova (SN) explosion. Massive stars lose their nucleosynthetic products to the Interstellar Medium (ISM) via strong winds (during the MS and WR phase) and during the SN explosion. Interestingly, while $^{60}$Fe is released only during the SN explosion, $^{26}$Al is released during the MS, the WR{,} and the SN phases \\citep{Limongi2006, Palacios2005}.   In the classical SN model, first proposed by \\citet{Cameron1977}, just after the discovery of $^{26}$Al \\citep{Lee1976}, $^{26}$Al and $^{60}$Fe were delivered together by a single SN \\citep{Boss2010, Ouellette2010} into the nearby solar protoplanetary disk or prestellar core. The distance, $r$,  at which {an} SN needs to be in order to inject {an} SLR at the solar abundance into a phase (prestellar core or protoplanetary disk) having a linear size $r_{0}$ reads {as}\\citep{Cameron1995, Gounelle2008}:  $$ r= {r_0 \\over 2} \\sqrt{ {\\eta_{\\rm SN} Y_{\\rm SN}\\over M_{\\rm SLR}} e^{-\\Delta/\\tau}, } \\eqno (1) $$ where $\\eta_{\\rm SN}$ is the mixing efficiency of the SN ejecta with the receiving phase, $Y_{\\rm SN}$ is the SN yield of the considered SLR, $M_{\\rm SLR}$ is the solar system mass of the SLR, $\\Delta$ the time elapsed between the release of the radioactive element by the source and its incorporation in the receiving phase{,} and $\\tau$ the mean lifetime of the SLR under scrutiny.   Because the solar system abundance of  $^{26}$Al is far better constrained than that of $^{60}$Fe, we use the former SLR (C$_{\\sun}$ [$^{26}$Al] = 3.3 ppb) to calculate the maximum distance at which the SN has to be from either the disk or the core to deliver SLRs at the solar abundance. Using the SN yields of massive stars {with} M $\\le$ 60 M$_{\\sun}$ (Huss et al. 2009) calculated by Woosley \\& Heger (2007), we  obtain r $\\le$ 0.4 pc for a disk of mass 0.013 M$_{\\sun}$ \\citep{Hayashi1985,Ouellette2005} and size r$_{0}$ = 100 AU with an injection efficiency of 0.7 \\citep{Ouellette2010}, and r $\\le$ 0.6 pc for a core of mass 2 M$_{\\sun}$ and size r$_{0}$~=~0.058 pc with an injection efficiency of 0.02 \\citep{Boss2010b}. The receiving phases parameters (mass and size) correspond to {the} observed properties of disks and cores adopted by the tenants of the single SN model \\citep{Ouellette2005,Boss2010b}, while the injection efficiencies are estimated by the same authors \\citep{Ouellette2010,Boss2010b}.  As  we conservatively assumed $\\Delta$ = 0, the calculated distances are upper limits of the maximum distances. In other words, if we applied a decay interval of  $\\Delta \\gtrsim$ 1 Myr as required by all SN models (Huss et al. 2009), it would imply that the receiving phases (disk or core) lie at a few tenths of a parsec from the SN at most to receive $^{26}$Al at the solar abundance.      \\begin{figure} \\includegraphics[width=0.38\\textwidth]{Figsketch-1.eps} \\caption{Sketch of the model described in the text (see Sect. 2). Star generations \\#1, \\#2, and \\#3 are respectively in red, blue, and green. In panel (b), SNe remnants are shown with a dark contour. The reddish background symbolizes $^{60}$Fe delivered by the SNe belonging to the first star generation. In panel (c), the purple color of the shell symbolizes the combination of $^{26}$Al delivered by the single massive star wind from generation \\#2 and the $^{60}$Fe delivered by generation \\#1 stars. Our Sun, born in the circumstellar shell together with a few hundreds fellow low-mass stars (generation \\#3), is shown in yellow.} \\label{figsketch} \\end{figure}  It is very unlikely, if not impossible, {however,} to find a protoplanetary disk or a dense core that close to {an} SN. Before they explode as SNe, massive stars carve large ionized regions in the ISM (called HII regions) where the gas density is too low and temperature too high for star formation to take place \\citep{Bally2008}. Observations show that even around a massive star that still needs to evolve for 2 Myr before it explodes as a SN, disks and cores are found several parsecs away \\citep{Hartmann2005}, too far to receive $^{26}$Al and $^{60}$Fe at the solar abundance. In addition, because SNe ejecta are vastly enriched in $^{60}$Fe relative to $^{26}$Al and their respective solar abundances (Woosley \\& Weaver 2007), all models relying on SN injection lead to a $^{26}$Al/$^{60}$Fe ratio {that is} far lower than the initial solar ratio of 8.2, unless special conditions are adopted \\citep{Desch2011}.   Following the pioneering work of \\citet{Arnould1997}, WR stars winds {have} recently { been} reconsidered as a specific source for $^{26}$Al {alone} \\citep{Arnould2006, Gaidos2009, Tati2010}. In the model of Gaidos et al. (2009), which considers injection of $^{26}$Al at the molecular cloud scale, very specific conditions (such as a precise timing between the formation of massive stars and the Sun or stellar clusters with an extremely large number of stars) are needed. In the stimulating model of \\citet{Tati2010}, $^{26}$Al is delivered into a bow-shock-produced shell by a single runaway massive star moving with a velocity $\\ge$ 20 km/s in a dense (n~$\\sim$~100 cm$^{-3}$) star-forming region. % At such a velocity ($\\sim$ 20 pc/Myr), the runaway WR star considered by \\citet{Tati2010} would escape any dense star-forming region with size $\\ge$ 40 pc within 2 Myr, preventing the {collection} of dense gas well before the entry into the WR phase. In addition none of these models is generic (i.e. they fail to offer a common explanation for $^{60}$Fe and $^{26}$Al), nor do they constrain the solar system{'s} genealogy. Finally, because of the rarity of WR stars \\citep{Crowther2007}, such models somehow require a special explanation for the formation of the solar system.  The goal of the present work is to identify a coherent model {that} accounts for the presence of $^{26}$Al and $^{60}$Fe in the nascent solar system and which is in line with the most recent astronomical observations of star-forming regions (Sect. 2). The proposed origin for $^{26}$Al (Sects. 2 and 3) is entirely original and relies on new rotating models of massive stars. The proposed explanation for $^{60}$Fe (Sect. 4) is an update of the Supernova Propagation \\& Cloud Enrichment (SPACE) model elaborated by Gounelle et al. (2009). Combining these two results, a generic explanation is offered for the presence of SLRs having a diversity of mean lifetimes in the early solar system (Sect. 5). In Sect. 6, we discuss our results, focusing on the newly established solar system genealogy. ",
        "conclusions": "We have elucidated the origin of SLRs in the early solar system by developing a new model for $^{26}$Al {that} relies on a physical mechanism (collect + injection, collapse) similar in essence to the one presented recently for $^{60}$Fe \\citep{Gounelle2009}. This new mechanism occurs naturally within a common mode of star formation, namely that of  sequential star formation within a {giant molecular cloud} \\citep{Hennebelle2009}. Within a few milligrams of meteorites, SLRs therefore record physical mechanisms observed in the sky {on} scales varying from hundreds down to 1 pc.  The identified sequence of events establishes the genealogy of the solar system. Our Sun is the {great}-grandson of a star complex (generation 0) containing 10s of thousands of stars, the grandson of a large GMC core (generation 1) containing a few thousand stars, {and} the son of a massive ($\\gtrsim$32 M$_{\\sun}$) star belonging to a cluster (1000-2000 stars with a preferred value of 1200 stars, generation \\#2) born later within the same GMC. Assuming 30  \\% star formation efficiency \\citep{Lada2003} and with an average stellar mass of M$_\\star$ = 0.5 M$_{\\sun}$, our Sun was born together with $\\sim$ six hundred fellow stars in its natal 1000 M$_{\\sun}$ shell.  Formation of the Sun in a relatively small clusteris in agreement with dynamical requirements, i.e. stability of planetary orbits, existence of the Kuiper belt object Sedna{,} and formation of the Oort cloud \\citep{Adams2010, Brasser2011}.   Given the size of its cluster at birth (six hundred stars), it is not {in}conceivable that our Sun was coeval to a massive star. {Since t}he average number of massive stars in a cluster of size N {is} f$_{\\rm SN}$ $\\times$ N, with f$_{\\rm SN}$ = 2.3 $\\times$ 10$^{-3}$, clusters with six hundred stars might contain between 1 or 2 massive stars, most probably B stars, leaving  room for disk photo evaporation as suggested by \\citet{Throop2005}.  It is expected that the few hundred stars born together with our Sun will share the same chemical and isotopic properties. These stars were true twins of our Sun. Given that stellar clusters dissipate on timescales of 100 Myr \\citep{Adams2001}, that revolution timescales around the Galactic center are {on} the order of 200 Myr{,} and that stars can radially migrate over a few kpc in the Galaxy \\citep{Roskar2008}, these fellow stars are now in totally unrelated places. Because the $^{26}$Al enrichment mechanism we have identified is generic, many other stars are, however, expected to contain $^{26}$Al at a level close to that of our solar system."
    },
    "1208/1208.1388_arXiv.txt": {
        "abstract": "We reconstruct an $f(R)$ gravity model that gives rise to the particular $\\Lambda$CDM background evolution of the universe. We find well-defined, real-valued analytical forms for the $f(R)$ model to describe the universe both in the early epoch from the radiation to matter dominated eras and the late time acceleration period.  We further examine the viability of the derived $f(R)$ model and find that it is viable to describe the evolution of the universe in the past and there does not exist the future singularity in the Lagrangian. ",
        "introduction": "Cosmological observations from supernovae\\cite{1}, BAO~\\cite{BAOmeasurement}and CMB \\cite{WMAP} indicate that our universe is undergoing a phase of accelerated expansion. Understanding the nature of the cosmic acceleration is one of the biggest questions in modern physics. This acceleration is believed to be driven by a so called dark energy (DE) in the framework of Einstein's general relativity. The simplest explanation of such DE is the cosmological constant. However, the measured value of the cosmological constant is far below the prediction of any sensible quantum field theories and furthermore the cosmological constant leads inevitably to the coincidence problem, namely why the energy densities of matter and the vacuum are of the same order today(see \\cite{sean} for review).  Alternatively, the acceleration can be explained by modifying the gravity theory. The theory of general relativity might not be ultimately correct on cosmological scales. One of the simplest attempts is called $f(R)$ gravity, in which the scalar curvature in the Lagrangian density of Einstein's gravity is replaced by an arbitrary function of $R$. However the complexity of the field equations makes it difficult to obtain a viable $f(R)$ model to satisfy both cosmological and local gravity constraints\\cite{fr}. Recently, there appeared a useful approach to reconstruct the $f(R)$ model by inverting the observed expansion history of the universe to deduce what class of $f(R)$ theories give rise to the particular cosmological evolution  \\cite{Song}\\cite{Pogosian}\\cite{Lombriser}\\cite{dobado}\\cite{Dunsby}\\cite{Nojiri}\\cite{solution}. Some analytical forms for $f(R)$ gravity that admit the $\\Lambda$CDM expansion history in the background spacetime were constructed \\cite{dobado} \\cite{Dunsby}\\cite{Nojiri}. However, it was argued that only a simple real-valued expression of $f(R)$ model in the Lagrangian could admit an exact $\\Lambda$CDM expansion history\\cite{Dunsby}.  In this paper, we will further study this problem. We will perform a number of explicit reconstructions which lead to a number of interesting results. We will show that we can derive a well-defined real-valued analytical $f(R)$ in terms of the hypergeometric functions to admit an exact $\\Lambda$CDM expansion history. We will explicitly show that the $f(R)$ gravity not only can admit an exact $\\Lambda$CDM expansion in the recent epoch of the universe but also can admit an exact $\\Lambda$CDM expansion in the early time of the universe. We will also discuss the physical boundary conditions for these constructions.   This paper is organized as follows: In section~\\ref{RM}, we review the background dynamics of the universe in the $f(R)$ gravity and present the well-defined, real-valued analytical $f(R)$ forms that can exactly reproduce the same background expansion as that of the $\\Lambda$CDM model from the radiation dominated epoch to the matter dominated epoch. In section~\\ref{MA}, we present the explicit form for a $f(R)$ model that can mimic the evolution of the universe from the matter dominated epoch to the late time acceleration. We will also discuss the physical boundary conditions and viability for these models. In section~\\ref{conclusions}, we will summarize and conclude this work. ",
        "conclusions": "} In summary, in this work we have constructed an $f(R)$ gravity model that mimics the $\\Lambda$CDM universe expansion in both the early and late epochs. We found that there exists a real-valued function for the Ricci scalar in terms of hypergeometric functions which can give rise to the particular cosmological evolution of the $\\Lambda$CDM model. Although the constructed $f(R)$ model has weakness in describing the future expansion of the universe in the Einstein frame, in the Jordan frame it is viable to describe the past evolution of the universe and it does not have the future singularity in the Lagrangian.  The fact that the Lagrangian is well-defined demonstrates that this family of $f(R)$ models are no longer just simply phenomenological models, but the field equations instead can be deduced from the principle of least action. Furthermore, when $w_1 \\neq 0$, the constant $\\Lambda$ in Eq.(\\ref{viable2}) cannot be explained as the energy density of the vacuum and the model does not suffer the cosmological constant problem even though it has the same background expansion of the universe as the $\\Lambda$CDM model. For the background evolution of the universe, we cannot distinguish these $f(R)$ models from general relativity. However, we can distinguish them at the cosmological perturbation level since the $f(R)$ gravity introduces an extra scalar degree of freedom which has significant impact on perturbation equations. The constraints from observations at linear perturbation level for the $f(R)$ model have been presented in our companion work ~\\cite{He}. However, it would be more interesting to investigate the nonlinear behavior in f(R) gravity using N-body simulations. The analytical functional form of $f(R)$ plays a vital role in N-body simulations, which is a subject of our future work.     \\emph{Acknowledgments J.H.He acknowledges the Financial support of MIUR through PRIN 2008 and ASI through contract Euclid-NIS I/039/10/0. The work of B.Wang was partially supported by NNSF of China under grant 10878001 and the National Basic Research Program of China under grant 2010CB833000.}"
    },
    "1208/1208.3491_arXiv.txt": {
        "abstract": "We describe the implementation of a search for gravitational waves from compact binary coalescences in LIGO and Virgo data. This all-sky, all-time, multi-detector search for binary coalescence has been used to search data taken in recent LIGO and Virgo runs.  The search is built around a matched filter analysis of the data, augmented by numerous signal consistency tests designed to distinguish artifacts of non-Gaussian detector noise from potential detections.   We demonstrate the search performance using Gaussian noise and data from the fifth LIGO science run and demonstrate that the signal consistency tests are capable of mitigating the effect of non-Gaussian noise and providing a sensitivity comparable to that achieved in Gaussian noise. ",
        "introduction": "\\label{sec:intro}  Coalescing binaries of compact objects such as \\acp{NS} and stellar-mass \\acp{BH} are promising \\ac{GW} sources for ground-based, kilometer-scale interferometric detectors such as LIGO \\cite{Abbott:2007kv}, Virgo \\cite{Accadia:2012zz}, and GEO600 \\cite{Grote:2008}, which are sensitive to waves of frequencies between tens and thousands of Hertz. Numerous searches for these signals were performed on data from the six LIGO and GEO science runs (S1--S6) and from the four Virgo science runs (VSR1--4) \\cite{Abbott:2003pj,LIGOS2iul,LIGOS2bbh,LIGOS2macho,ligotama, S3_BCVSpin,LIGOS3S4all,Collaboration:2009tt,Abbott:2009qj,S5LowMassLV,Collaboration:S6CBClowmass}.  Over time, the software developed to run these searches and evaluate the significance of results evolved into a sophisticated pipeline, known as \\ihope. An early version of the pipeline was described in \\cite{brown-2005-22}. In this paper, we describe the \\ihope\\ pipeline in detail and we characterize its detection performance by comparing the analysis of a month of real data with the analysis of an equivalent length of simulated data with Gaussian stationary noise.  \\acfp{CBC} consist of three dynamical phases: a gradual \\emph{inspiral}, which is described accurately by the post-Newtonian approximation to the Einstein equations \\cite{Blanchet:2002av}; a nonlinear \\emph{merger}, which can be modeled with numerical simulations (see \\cite{Centrella:2010,Hannam:2009rd,2011arXiv1107.2819S} for recent reviews); and the final ringdown of the merged object to a quiescent state \\cite{Berti:2007gv}. For the lighter NS--NS systems, only the inspiral lies within the band of detector sensitivity. Since \\ac{CBC} waveforms are well modeled, it is natural to search for them by matched-filtering the data with banks of theoretical \\emph{template} waveforms \\cite{wainstein:1962}.  The most general \\ac{CBC} waveform is described by seventeen parameters, which include the masses and intrinsic spins of the binary components, as well as the location, orientation, and orbital elements of the binary. It is not feasible to perform a search by placing templates across such a high-dimensional parameter space. However, it is astrophysically reasonable to neglect orbital eccentricity \\cite{Cokelaer:2009hj,Brown:2009ng}; furthermore, \\ac{CBC} waveforms that omit the effects of spins have been shown to have acceptable phase overlaps with spinning-binary waveforms, and are therefore suitable for the purpose of detecting \\acp{CBC}, if not to estimate their parameters accurately \\cite{VanDenBroeck:2009gd}.  Thus, \\ac{CBC} searches so far have relied on nonspinning waveforms that are parameterized only by the component masses, by the location and orientation of the binary, by the initial orbital phase, and by the time of coalescence. Among these parameters, the masses determine the intrinsic phasing of the waveforms, while the others affect only the relative amplitudes, phases, and timing observed at multiple detector sites \\cite{Allen:2005fk}. It follows that templates need to be placed only across the two-dimensional parameter space spanned by the masses \\cite{Allen:2005fk}. Even so, past \\ac{CBC} searches have required many thousands of templates to cover their target ranges of masses. (We note that \\ihope\\ could be extended easily to nonprecessing binaries with aligned spins. However, more general precessing waveforms would prove more difficult, as discussed in \\cite{PhysRevD.49.6274,Apostolatos:1995,BuonannoChenVallisneri:2003b,Pan:2003qt}.)  In the context of stationary Gaussian noise, matched-filtering would directly yield the most statistically significant detection candidates. In practice, environmental and instrumental disturbances cause non-Gaussian noise transients (\\emph{glitches}) in the data. Searches must distinguish between the candidates, or \\textit{triggers}, resulting from glitches and those resulting from true \\acp{GW}. The techniques developed for this challenging task include \\emph{coincidence} (signals must be observed in two or more detectors with consistent mass parameters and times of arrival), \\emph{signal-consistency} tests (which quantify how much a signal's amplitude and frequency evolution is consistent with theoretical waveforms \\cite{Allen:2004}), and \\emph{data quality vetoes} (which identify time periods when the detector glitch rate is elevated). We describe these in detail later.  The \\emph{statistical significance} after the consistency tests have been applied is then quantified by computing the \\ac{FAP} or \\ac{FAR} of each candidate; we define both below. For this, the background of noise-induced candidates is estimated by performing \\emph{time shifts}, whereby the coincidence and consistency tests are run after imposing relative time offsets on the data from different detectors. Any consistent candidate found in this way must be due to noise; furthermore, if the noise of different detectors is uncorrelated, the resulting background rate is representative of the rate at zero shift.  The \\emph{sensitivity} of the search to \\ac{CBC} waves is estimated by adding simulated signals (\\emph{injections}) to the detector data, and verifying which are detected by the pipeline. With this diagnostic we can tune the search to a specific class of signals (e.g., a region in the mass plane), and we can give an astrophysical interpretation, such as an upper limit on \\ac{CBC} rates \\cite{Fairhurst:2007qj}, to completed searches.  As discussed below, commissioning a \\ac{GW} search with the \\ihope\\ pipeline requires a number of parameter tunings, which include the handling of coincidences, the signal-consistency tests, and the final ranking of triggers.  To avoid biasing the results, \\ihope\\ permits a blind analysis:  the results of the non-time-shifted analysis can be sequestered, and tuning performed using only the injections and time-shifted results.  Later, with the parameter tunings frozen, the non-time-shifted results can be unblinded to reveal the candidate GW events.  This paper is organized as follows. In Sec.\\ \\ref{sec:coinc_search} we provide a brief overview of the \\ihope\\ pipeline, and describe its first few stages (data conditioning, template placement, filtering, coincidence), which would be sufficient to implement a search in Gaussian noise but not, as we show, in real detector data. In Sec.\\ \\ref{sec:nongauss} we describe the various techniques that have been developed to eliminate the majority of background triggers due to non-Gaussian noise. In Sec.\\ \\ref{sec:interpretation} we describe how the \\ihope\\ results are used to make astrophysical statements about the presence or absence of signals in the data, and to put constraints on \\ac{CBC} event rates. Last, in Sec.\\ \\ref{sec:discussion} we discuss ways in which the analysis can be enhanced to improve sensitivity, reduce latency, and find use in the advanced-detector era.  Throughout this paper we show representative \\ihope\\ output, taken from a search of one month of LIGO data from the S5 run (the third month in \\cite{Abbott:2009qj}), when all three LIGO detectors (but not Virgo) were operational. The search focused on low-mass \\ac{CBC} signals with component masses $> 1 \\, M_{\\odot}$ and total mass $< 25 \\, M_{\\odot}$. For comparison, we also run the same search on Gaussian noise generated at the design sensitivity of the \\ac{LIGO} detectors (using the same data times as the real data). Where we perform \\ac{GW}-signal injections (see Sec.\\ \\ref{ssec:injections}), we adopt a population of binary-neutron-star inspirals, uniformly distributed in distance, coalescence time, sky position and orientation angles. ",
        "conclusions": "\\label{sec:discussion}   In this paper we have given a detailed description of the \\ihope\\ software pipeline, developed to search for \\acp{GW} from \\ac{CBC} events in LIGO and Virgo data, and we have provided several examples of its performance on a sample stretch of data from the LIGO S5 run. The pipeline is based on a matched-filtering engine augmented by a substantial number of additional modules that implement coincidence, signal-consistency tests, data-quality cuts, tunable ranking statistics, background estimation by time shifts, and sensitivity evaluation by injections. Indeed, with the \\ihope\\ pipeline we can run analyses that go all the way from detector strain data to event significance and upper limits on \\ac{CBC} rates.  The pipeline was developed over a number of years, from the early versions used in LIGO's S2 BNS search to its mature incarnation used in the analysis of S6 and VSR3 data. One of the major successes of the \\ihope\\ pipeline was the mitigation of spurious triggers from non-Gaussian noise transients, to such an extent that the overall volume sensitivity is reduced by less than 20\\%  compared to what would be possible if noise was Gaussian.  Nevertheless, there are still significant improvements that can and must be made to \\ac{CBC} searches if we are to meet the challenges posed by analyzing the data of advanced detectors. In the following paragraphs, we briefly discuss some of these improvements and challenges.  \\paragraph*{Coherent analysis.} As discussed above, the \\ihope\\ pipeline comes close to the sensitivity that would be achieved if noise was Gaussian, with the same PSD.  Therefore, while some improvement could be obtained by implementing more sophisticated signal-consistency tests and data-quality cuts, it will not be significant.  If three or more detectors are active, sensitivity \\emph{would} be improved in a \\emph{coherent} \\cite{FinnChernoff:1993, PaiDhurandharBose2001, HarryFairhurst:2011} (rather than coincident) analysis that filters the data from all operating detectors simultaneously, requiring consistency between the times of arrival and relative amplitudes of \\ac{GW} signals, as observed in each data stream. Such a search is challenging to implement because the data from the detectors must be combined differently for each sky position, significantly increasing computational cost.  Coherent searches \\emph{have} already been run for unmodeled burst-like transients \\cite{S5VSR1Burst}, and for \\ac{CBC} signals in coincidence with gamma-ray-burst observations \\cite{Briggs:2012ce}, but a full all-sky, all-time pipeline like  \\ihope\\ would require significantly more computation. A promising compromise may be a hierarchical search consisting of a first coincidence stage followed by the coherent analysis of candidates, although the estimation of background trigger rates would prove challenging as time shifts in a coherent analysis cannot be performed using only the recorded single detector triggers but require the full \\ac{SNR} time series.  \\paragraph*{Background estimation.} The first positive \\ac{GW} detection requires that we assign a very low false-alarm probability to a candidate trigger \\cite{Collaboration:S6CBClowmass}. In the \\ihope\\ pipeline, this would necessitate a large number of time shifts, thus negating the computational savings of splitting matched filtering between two stages, or a different method of background estimation \\cite{Dent:2012,Cannon:2012}.  Whichever the solution, it will need to be automated to identify signal candidates rapidly for possible astronomical follow up.  \\paragraph*{Event-rate estimation.} After the first detections, we will begin to quote event-rate estimates rather than upper limits. The loudest-event method can be used for this \\cite{Biswas:2007ni}, provided that the data are broken up so that much less than one gravitational wave signal is expected in each analyzed stretch. There are however other approaches \\cite{Messenger:2012} that should be considered for implementation.  \\paragraph*{Template length.} The sensitive band of advanced detectors will extend to lower frequencies (\\(\\unit{\\sim 10}{\\hertz}\\)) than their first-generation counterparts, greatly increasing the length and number of templates required in a matched-filtering search. Increasing computational resources may not be sufficient, so we are investigating alternative approaches to filtering \\cite{Marion:2004, Cannon:2010, Cannon:2011tb, Cannon:2011xk, Cannon:2011vi} and possibly the use of graphical processing units (GPUs).  \\paragraph*{Latency.} The latency of \\ac{CBC} searches (i.e., the ``wall-clock'' time necessary for search results to become available) has decreased over the course of successive science runs, but further progress is needed to perform prompt follow-up observations of \\ac{GW} candidate with conventional (electromagnetic) telescopes \\cite{Virgo:2011aa, Metzger:2011bv}. The target should be posting candidate triggers within minutes to hours of data taking, which was in fact achieved in the S6--VSR3 analysis with the MBTA pipeline \\cite{Marion:2004}.  \\paragraph*{Template accuracy.} While the templates currently used in \\ihope\\ are very accurate approximations to \\ac{BNS} signals, they could still be improved for the purpose of \\ac{NSBH} and \\ac{BBH} searches \\cite{BuonannoIyerOchsnerYiSathya2009}. It is straightforward to extend \\ihope\\ to include the effects of spin on the progress of inspiral (i.e., its \\emph{phasing}), but it is harder to include the orbital precession caused by spins and the resulting waveform modulations. The first extension would already improve sensitivity to \\ac{BBH} signals \\cite{Ajith:2009bn, Santamaria:2010}, but precessional effects are expected to be more significant for \\ac{NSBH} systems \\cite{Pan:2003qt, Ajith:2011hq}.  \\paragraph*{Parameter estimation.} Last, while \\ihope\\ effectively searches the entire template parameter space to identify candidate triggers, at the end of the pipeline the only information available about these are the estimated binary masses, arrival time, and effective distance. Dedicated follow-up analyses can provide much more detailed and reliable estimates of all parameters \\cite{Sluys:2008a, Sluys:2008b, Veitch:2010, Feroz:2009}, but \\ihope\\ itself could be modified to provide rough first-cut estimates."
    },
    "1208/1208.6082_arXiv.txt": {
        "abstract": "Recently, an interesting indication for a dark matter signal in the form of a narrow line, or maybe two lines and/or an internal bremsstrahlung feature, has been found in analyses of public data from the Fermi-LAT satellite detector. As recent analyses have also shown that there is little sign of extra contributions to continuum photons, it is natural to investigate leptophilic interacting massive particle (LIMP) models. We show that a model of radiatively generated neutrino masses may have the properties needed to explain the Fermi-LAT structure around 130 GeV. This model was proposed some 10 years ago, and predicted a clearly observable $\\gamma$-ray signal in the Fermi-LAT (then GLAST) detector. Here, we update and improve that analysis, and show as an example that a right-handed neutrino of mass 135 GeV should give rise to three conspicuous effects: a broad  internal bremsstrahlung bump with maximum around 120 GeV, a 2$\\gamma$ line around 135 GeV, and a $Z\\gamma$ line at 119.6 GeV (neglected in the previous work). These features  together give a good fit to the 130 GeV structure, given the present energy resolution of the Fermi-LAT data. An attractive feature of the model is that the particle physics properties are essentially fixed, once the relic density and the mass of the right-handed neutrino dark matter particle have been set. Puzzling features of the data at present are a slight displacement of the signal from the galactic center, and a needed boost factor of order $5-15$. This presents interesting challenges for numerical simulations including both baryons and dark matter on scales of 100 pc, and perhaps a need to go beyond the simplest  halo models. With upcoming experiments having better energy resolution, or with future Fermi-LAT data, the double-peak structure with a definite predicted ratio of the strengths of the two lines and the internal bremsstrahlung feature should be seen, if this model is correct. With the planned satellite GAMMA-400, a striking fingerprint of this dark matter candidate should then appear. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1208/1208.3202_arXiv.txt": {
        "abstract": "Absorption lines from the molecules OH$^+$, H$_2$O$^+$, and H$_3^+$ have been observed in a diffuse molecular cloud along a line of sight near W51 IRS2.  We present the first chemical analysis that combines the information provided by all three of these species.  Together, OH$^+$ and H$_2$O$^+$ are used to determine the molecular hydrogen fraction in the outskirts of the observed cloud, as well as the cosmic-ray ionization rate of atomic hydrogen.  H$_3^+$ is used to infer the cosmic-ray ionization rate of H$_2$ in the molecular interior of the cloud, which we find to be $\\zeta_2=(4.8\\pm3.4)\\times10^{-16}$~s$^{-1}$.  Combining the results from all three species we find an efficiency factor---defined as the ratio of the formation rate of OH$^+$ to the cosmic-ray ionization rate of H---of $\\epsilon=0.07\\pm0.04$, much lower than predicted by chemical models.  This is an important step in the future use of OH$^+$ and H$_2$O$^+$ on their own as tracers of the cosmic-ray ionization rate. ",
        "introduction": "\\label{section_intro}  \\setcounter{footnote}{6} \\renewcommand{\\thefootnote}{\\arabic{footnote}}  In the past decade, H$_3^+$ has widely become regarded as an excellent tracer of the cosmic-ray ionization rate in diffuse molecular clouds.  Surveys of H$_3^+$ in such clouds \\citep{indriolo2007,indriolo2012} have enabled us to find variations in the ionization rate between sight lines, and to build up the distribution function of cosmic-ray ionization rates in the nearby interstellar medium (ISM).  However, observations of H$_3^+$ are currently limited to background sources with $L$-band magnitudes brighter than about $L=7.5$~mag.  At this cutoff, OB stars are only feasible as background sources to distances of a few kpc, meaning that H$_3^+$ observations are primarily limited to the local spiral arm \\citep[observations toward the Galactic center, e.g., ][use dust-embedded objects]{goto2002,goto2008,goto2011,oka2005,geballe2010}.  An alternative method for inferring the ionization rate utilizes the chemistry associated with the formation and destruction of OH$^+$ and H$_2$O$^+$ \\citep{gerin2010,neufeld2010}, thought to be dependent primarily on hydrogen abstraction reactions with H$_2$ and dissociative recombination with electrons.  The HIFI instrument \\citep{degraauw2010} aboard {\\em Herschel} \\citep{pilbratt2010} has provided the first opportunity to observe both OH$^+$ and H$_2$O$^+$ with very high spectral resolution, thus allowing the use of these ions in constraining the cosmic-ray ionization rate.  Background sources bright enough for THz spectroscopy are widely distributed throughout the Galaxy, and targets from the PRISMAS\\footnote{PRobing InterStellar Molecules with Absorption line Studies} key programme range in distance from about 1~kpc to 12~kpc.  However, the ionization rate inferred from the oxygen chemistry is dependent upon an efficiency factor, $\\epsilon$, at which atomic hydrogen ionized by cosmic rays will eventually be converted into OH$^+$.  In order to determine $\\epsilon$, we present observations of OH$^+$, H$_2$O$^+$, and H$_3^+$ in sight lines toward W51 and compare the ionization rates inferred separately from the hydrogen chemistry and oxygen chemistry.  \\subsection{Hydrogen Chemistry}  The interstellar chemistry of H$_3^+$ is rather simple, and the reactions surrounding this molecule are given in the top portion of Table \\ref{tbl_reactions}.  H$_3^+$ is formed in a two-step process, beginning with the ionization of H$_2$ by cosmic rays, and quickly followed by a reaction of H$_2^+$ with H$_2$.  Some H$_2^+$ is destroyed by dissociative recombination with electrons or by charge transfer to atomic hydrogen, but these reactions are generally slow compared to the $\\mathrm{H_2^+ + H_2}$ process.\\footnote{This is no longer true for the $\\mathrm{H_2^+ + H}$ reaction at low molecular fraction.}  Cosmic-ray ionization is the rate-limiting step in this process as it is many orders of magnitude slower than proton transfer from H$_2^+$ to H$_2$, and can be taken as the formation rate of H$_3^+$.  The primary destruction mechanisms for H$_3^+$ are dependent on the environment under consideration.  In diffuse molecular clouds, H$_3^+$ is predominantly destroyed via dissociative recombination with electrons.  In dense clouds, however, where the electron fraction is much lower, H$_3^+$ is destroyed by proton transfer to neutrals such as CO and O. {%    \\subsection{Oxygen Chemistry}  Reactions involved in the chemistry surrounding OH$^+$ and H$_2$O$^+$ are presented in the bottom portion of Table \\ref{tbl_reactions}.  The formation of OH$^+$ begins with the ionization of atomic hydrogen by cosmic rays.  This is followed by endothermic charge transfer to oxygen to form O$^+$---a process highly dependent upon the relative populations in the fine structure levels of atomic oxygen \\citep{stancil1999}---and hydrogen abstraction from H$_2$ to form OH$^+$.  OH$^+$ is either destroyed by further hydrogen abstraction to form H$_2$O$^+$, or by dissociative recombination with electrons.  The same is true for H$_2$O$^+$, but H$_3$O$^+$ is only destroyed by dissociative recombination with electrons.  A steady-state analysis of these reactions is employed in Section \\ref{section_analysis} in inferring the ionization rate of atomic hydrogen and molecular hydrogen fraction from OH$^+$ and H$_2$O$^+$ abundances.  An alternative means of forming OH$^+$ is the reaction of O with H$_3^+$.  This process requires a high molecular hydrogen fraction---such that H$_3^+$ is formed efficiently from cosmic-ray ionization of H$_2$---and a low electron fraction---such that H$_3^+$ is predominantly destroyed by proton transfer to O, forming OH$^+$.  As we will show in Section \\ref{section_analysis} that the OH$^+$ and H$_2$O$^+$ probed by our observations reside in gas with a low molecular hydrogen fraction, we omit this pathway from our analysis. ",
        "conclusions": "The efficiency factor we determine is much lower than that predicted by PDR models computed using the Meudon code \\citep{lepetit2006,goicoechea2007}, where  $0.5\\lesssim\\epsilon\\leq1.0$ \\citep[see discussion in][]{neufeld2010}.  In those models, the chain of reactions leading from H$^+$ to OH$^+$ is broken by recombination of H$^+$ or O$^+$ with electrons, both of which decrease $\\epsilon$.  Although recombination of O$^+$ is not important, recombination of H$^+$, while slow\\footnote{$k({\\rm H}^+|e^-)=3.5\\times10^{-12}(T/300)^{-0.75}$~cm$^{3}$~s$^{-1}$; UDFA06}, can compete with the O$^+$ + H$_2$ reaction at low molecular fraction where the $\\mathrm{O^+ + H\\rightarrow O + H^+}$ back-reaction dominates the reaction with H$_2$ that forms OH$^+$.  As a result, the reaction network cycles between H$^+$ and O$^+$, sometimes forming OH$^+$ and sometimes forming H.  Because of this mechanism, at $f_{\\rm H_2}=0.04$ the efficiency factor is about 0.5.  Another property of the Meudon models that can reduce $\\epsilon$ is the inclusion of state-specific rate coefficients for the H$^+$ + O reaction, and the relative populations in the fine-structure levels of $(^3{\\rm P}_J)$O.  The ground ($J=2$) level is likely the most populated state, and the H$^+$ + $(^3{\\rm P}_2)$O reaction is dramatically slower at low ($\\lesssim100$~K) temperatures than reactions involving $(^3{\\rm P}_1)$O or $(^3{\\rm P}_0)$O \\citep{stancil1999}.  If charge transfer to O$^+$ is inhibited, then H$^+$ has more time to recombine with electrons, thus decreasing the efficiency at which OH$^+$ forms.   Still, recombination of H$^+$ alone cannot account for the small value of $\\epsilon=0.07$ that we find.  In order to reach lower values of $\\epsilon$ under diffuse cloud conditions, something other than electrons must be removing H$^+$ from the gas phase.  One possible mechanism for doing just this is the neutralization of H$^+$ on small grains or PAH's \\citep{liszt2003}.  Recent modeling efforts \\citep{hollenbach2012} that account for grain/PAH neutralization find values of $\\epsilon\\sim0.1$--0.3, much closer to our observationally-derived value.  We can apply these low efficiency factors to update ionization rates inferred in previous studies of OH$^+$ and H$_2$O$^+$.  Rescaling the lower limit on the ionization rate reported in \\citet{gerin2010} to account for $\\epsilon=0.07$, we find $\\zeta_{\\rm H}>2.6\\times10^{-18}n({\\rm H})$~s$^{-1}$ for the line of sight toward W31C.  Doing the same for the range of ionization rates reported in \\citet{neufeld2010} toward W49N results in $8.6\\times10^{-16}$~s$^{-1}\\leq\\zeta_{\\rm H}\\leq17\\times10^{-16}$~s$^{-1}$.  These values are high compared to the distribution of ionization rates found using H$_3^+$ \\citep[$\\zeta_{\\rm H}=(2.3_{-2.0}^{+3.4})\\times10^{-16}$~s$^{-1}$, converted from the mean value of $\\zeta_2$ in][]{indriolo2012}.  However, this scaling procedure is highly uncertain at present, and must be improved by determining $\\epsilon$ in more cloud components.  It is also possible that our analysis underestimates $\\epsilon$ given uncertainty in the rate coefficient for dissociative recombination of OH$^+$ with electrons.  As shown in Table \\ref{tbl_reactions}, this process is 10 times slower than all of the other dissociative recombination reactions.  However, it has not been measured at temperatures relevant to diffuse interstellar clouds using vibrationally cold OH$^+$ molecules.  Under such conditions, it is possible that resonance structure in the low-energy collision cross section may increase the low-temperature rate of this reaction.  Indeed, some resonance structure has been observed \\citep{amitay1996}, but that experiment did not determine the cross section on an absolute scale.  Future measurements of the OH$^+$ + $e^-$ dissociative recombination cross section using storage ring facilities are urgently needed.\\footnote{If $k({\\rm OH}^+|e^-)$ is 10 times larger than the value adopted in this study (see Table \\ref{tbl_reactions}), then the value of $\\epsilon$ required to bring the cosmic-ray ionization rates inferred from H$_3^+$ and OH$^+$ into agreement increases from 0.07 to 0.23}  Additionally, it would be advantageous to employ a chemical model that is more complete than the analytical expressions used herein.  Several species have been observed in the diffuse molecular cloud toward W51, all of which can be used in constraining the ambient physical conditions.  However, some species---e.g., H \\citep{koo1997}, CH$^+$ \\citep{falgarone2010}, OH$^+$, H$_2$O$^+$---are thought to reside in primarily atomic gas, while others---e.g., CH \\citep{gerin2010ch}, HF \\citep{sonnentrucker2010}, H$_2$O \\citep{neufeld2002,sonnentrucker2010}, HCO$^+$, HCN \\citep{godard2010}, CO---prefer molecular environments.  It seems apparent then, that even with similar velocity profiles, not all of the observed species are spatially co-located.  Instead, these atoms and molecules are likely probing different portions of a cloud complex, including diffuse atomic outer layers, and a diffuse molecular interior.  Any chemical model attempting to reproduce the observed abundances along this sight line must account for these effects.  Our analysis of OH$^+$ and H$_2$O$^+$ shows that both species must reside in gas of low molecular hydrogen fraction ($f_{\\rm H_2}=0.04$) in order to explain the observed abundance ratio between the two species.  H$_3^+$, however, is not efficiently formed in gas with low molecular fractions where the H$_2^+$ + H reaction competes with the H$_2^+$ + H$_2$ reaction, meaning that H$_3^+$ must primarily reside in regions of higher molecular hydrogen fraction.  As such, it seems likely that there is little overlap in the gas probed by OH$^+$ and H$_2$O$^+$, and that probed by H$_3^+$ in diffuse molecular clouds (this changes for gas with a low electron fraction where the H$_3^+$ + O reaction becomes important)."
    },
    "1208/1208.6427_arXiv.txt": {
        "abstract": "s{ Several projects in radioastronomy plan to use large static cylindrical reflectors with an extended lobe sampling a sector of the rotating sky. This study provides the exact mathematical expression of the transit time of a celestial object within the acceptance lobe of such a cylindrical device. The mathematical approach, based on the stereographic projection, allows one to study the optimisation of the position and orientation of the radio-reflector, and should provide exact coefficients for the spatial Fourier Transform of the radio signal along the cylinder axis. \\\\ {\\bf Keywords}: Instrumentation: interferometers -- Cosmology: large-scale structure of Universe -- dark energy -- Radio lines: galaxies } ",
        "introduction": "Several baryonic oscillation (BAO) radio projects \\cite{1,2,3} plan to operate a series of parallel static reflectors of large parabolic cylinder shape, to map the $21 cm$ HI emission line. The sky will transit over the acceptance lobe, which is defined by an angular sector of aperture $\\Delta$ centered around a vertical plane (see Fig. \\ref{secteur}). \\begin{figure}[h] \\parbox{7cm}{ \\includegraphics[width=7cm]{dessin1.pdf} } \\hspace{0.5cm} \\parbox{6cm}{ \\caption[]{\\it The reflector and its field of view, defined as the angular sector $\\Delta$ that is focalised within the antenna's acceptance. } } \\label{secteur} \\end{figure}  The projection on the sky of this angular sector can be seen on Fig. \\ref{vision3D}. In this paper, I produce the exact calculation of the transit time of a celestial object, as a function of its declination. In section 5, I use the results of the calculation to compare the performances of various radio-telescope latitudes (France, Morocco, South Africa, equator) and configurations (orientation). In particular, we show that the North-South orientation usually considered may not be the optimal one for high-z BAO studies that need large exposure times, but not necessarilly the largest possible field of view.  Another possible use of the exact expression for the transit time is the production of exact coefficients for Fourier Transform calculations along the cylinder axis.  \\begin{figure}[h] \\centering \\vbox{ \\includegraphics[width=12cm]{dessin2.pdf} } \\caption[]{\\it The celestial sphere with the projected position of the observer $M_0$ (latitude $\\lambda$), the projected orientation of the reflector (A) and the projected portion of detectable sky (detection lobe), defined as the angular sector $\\Delta$ of axis $P_0 P'_0$ (in grey), where $P_0$ and $P'_0$ are the projections of the reflector's axis. $\\theta/2\\pi$ is the fraction of the sideral day that an object of declination $\\delta$ will spend within the detection lobe. } \\label{vision3D} \\end{figure} ",
        "conclusions": "The purpose of this study is to provide the exact expression of the transit time of a given celectial object within the lobe of a cylindrical reflector. Each particular case has been examined and the results have been used to analyse different telescope configurations. It is clear from this study that the groups planning to use a static setup of cylinders should seriously consider the orientation as a degree of freedom to favour either the largest field coverage with the shortest mean transit time (North-South orientation), or a smaller field coverage, but allowing a deeper survey (East-West orientation)."
    }
}